{
    "0910/0910.0878_arXiv.txt": {
        "abstract": "We compare the accuracy of various methods for determining the transfer of the diffuse Lyman continuum in H{\\sc\\,ii} regions, by comparing them with a high-resolution discrete-ordinate integration. We use these results to suggest how, in multidimensional dynamical simulations, the diffuse field may be treated with acceptable accuracy without requiring detailed transport solutions.  The angular distribution of the diffuse field derived from the numerical integration provides insight into the likely effects of the diffuse field for various material distributions. ",
        "introduction": "In this paper, we model the diffuse field structure of astrophysical H{\\sc\\,ii} regions.  Recombinations within these regions emit radiation in the various hydrogen line spectra and Balmer and higher continua which is an observed characteristic of them.  Emission in the Lyman continuum, however, is energetic enough to ionize other hydrogen atoms, and so is trapped within the nebula.  This radiation field is believed to have a significant fraction of the energy density of the direct ionizing continuum in some parts of the H{\\sc\\,ii} region, and so it is important to model it correctly.  This is particularly relevant when modelling the complex dynamics of the nebulae, or the emission from the internal features such as the tails of cometary globules in the Helix planetary nebula \\citep{odell07}.  \\cite{ritze05} has suggested that diffuse fields may be particularly important at the edges of H{\\sc\\,ii} regions, in some regimes.  This would seem to suggest that the diffuse field may have a more significant impact on their overall evolution than previously assumed. However, Ritzerveld uses a simple outward-only treatment of the diffuse radiation transport, and also assumes that the absorption coefficient is comparable for diffuse and direct radiation fields, although it is acknowledged that in reality the direct photons are likely to have a harder spectrum and will thus be significantly more penetrating.  To investigate the validity of these conclusions, in this paper we apply detailed discrete-ordinate angular integration to investigate the validity of several approximate numerical transport schemes.  To simplify the problem, we assume a pure-hydrogen nebula, without dust, and use a simple two-frequency approximation to the radiation flow.  With our more detailed modelling, we find that the diffuse field can indeed dominate for the situations Ritzerveld describes.  However, where the diffuse field dominates, it will typically also be outwardly-beamed and therefore be indistinguishable from the direct field for most purposes.  For most astrophysically relevant conditions, the usual on-the-spot approximation is shown to give accurate results through most of a spherical nebula, except for a region close to the star where it {\\it overestimates}\\/ the diffuse field.  We also assess the accuracy of the Eddington diffusion approximation for the diffuse field transfer, which may be a useful means of modelling diffuse transport effects in multidimensional simulations.  In the present paper, we neglect the effects of dust and heavy elements in the H{\\sc\\,ii} region.  This is a reasonable assumption for the case of cosmological H{\\sc\\,ii} regions; however, there is observational evidence for the importance of dust absorption within the ionized gas in galactic H{\\sc\\,ii} regions \\citep{cesar00,robbe05}.  We briefly discuss the impact dust might have on our results. ",
        "conclusions": "We present detailed calculations of radiation transfer within a photoionized nebula, for a simplified physical model.  We find that the diffuse field is indeed enhanced in the outer parts of such a nebula if there are steep density gradients in the H{\\sc\\,ii} region, particularly if the absorption cross section for diffuse and direct photons is comparable.  These circumstances are not likely to be typical.  We also find that when the diffuse radiation is relatively strong, it is also strongly beamed in the radial direction, and so the dynamical effect of the radiation field will in any case be similar to the direct illumination.  We discuss the possibility of using a diffusion approximation for the diffuse radiation field in two- and three-dimensional radiation hydrodynamic calculations of H{\\sc\\,ii} regions.  This may allow the diffuse field effects to be calculated to reasonable accuracy without requiring a full radiation transfer calculation.  This paper has not considered in detail a number of additional processes which may effect the radiation field within the nebulae, such as the spectral hardening of radiation close to the ionization front, dust absorption and scattering, and the effects of helium and other heavy elements.  This is left for future work."
    },
    "0910/0910.2799_arXiv.txt": {
        "abstract": "{We report the results of a high-energy multi-instrumental campaign with \\integ, \\xte, and \\swift\\ of the recently discovered \\integ\\ source \\igr. The \\swift/XRT data allow us to refine the position of the source to RA$_{\\mathrm{J}2000}$= 19$^h$ 29$^m$ 55.9$^s$ Dec$_{\\mathrm{J}2000}$=+18$^\\circ$ 18\\arcmin\\ 38.4\\arcsec\\ ($\\pm 3.5$\\arcsec), which in turn permits us to identify a candidate infrared counterpart. The \\swift\\ and \\xte\\ spectra are well fitted with absorbed power laws with hard ($\\Gamma \\sim 1$) photon indices. During the longest \\swift\\ observation, we obtained evidence of absorption in true excess to the Galactic value, which may indicate some intrinsic absorption in this source. We detected a strong (P=40\\%) pulsation at $12.43781 (\\pm0.00003)$~s that we interpret as the spin period of a pulsar. All these results, coupled with the possible 117 day orbital period, point to \\igr\\ being an HMXB with a Be companion star. However, while the long-term \\integ/IBIS/ISGRI  18--40 keV light curve shows that the source spends most of its time in an undetectable state, we detect occurrences of short ($\\sim 2000-3000$~s) and intense flares that are more typical of supergiant  fast X-ray transients.  We therefore cannot make firm conclusions on the type of system, and we discuss the possible implications of \\igr\\ being an SFXT. } ",
        "introduction": "\\indent The INTErnational Gamma-Ray Astrophysics Laboratory (\\integ) has permitted a large number of new sources to be discovered. Amongst the $\\sim250$ new sources\\footnote{see http://isdc.unige.ch/$\\sim$rodrigue/html/igrsources.html for an up to date list of all sources and their properties.}, \\integ\\ has unveiled two new or poorly-known types of high mass X-ray binaries (HMXB): the very absorbed supergiant HMXBs, and the supergiant fast X-ray transients (SFXT). While all types of HMXBs (including those hosting a Be star, Be-HMXBs) are powered by accretion of material by a compact object, understanding the nature of a system is very important for studying the evolutionary paths of source populations, and more globally, the evolution of the Galaxy in terms of its source content.\\\\ \\indent \\igr\\ was discovered by \\citet{2009ATel.1997....1T} who reported activity of this source seen with \\integ\\ during the monitoring campaign of  \\grs\\ \\citep[e.g.][]{rodrigue08_1915a}. Archival \\swift\\ data of a source named \\object{Swift J1929.8+1818} allowed us to give a refined X-ray position of RA$_{\\mathrm{J}2000}$= 19$^h$ 29$^m$ 55.9$^s$ and Dec$_{\\mathrm{J}2000}$=+18$^\\circ$ 18\\arcmin\\ 39\\arcsec\\ ($\\pm 3.5$\\arcsec) which we used to locate a possible infrared counterpart \\citep{2009ATel.1998....1R}. We also suggested that the \\swift\\ and \\integ\\ sources are the same and that activity from \\igr\\ had been seen with \\swift\\ in the past.  The temporal analysis of the XRT light curve showed a possible pulsation at 12.4~s \\citep{2009ATel.1998....1R}. Analysis of \\swift/BAT archival data revealed that the source had been detected on previous occasions with a periodicity at 117.2 days \\citep{2009ATel.2008....1C} that we interpreted as the orbital period of the system. \\\\ \\indent Soon after the discovery of the source with \\integ, we triggered our accepted \\xte\\ and \\swift\\ programmes for follow-up observations of new \\integ\\ sources (PI Rodriguez). A preliminary analysis of the real-time \\xte\\ data allowed us to confirm the pulsations in the signal from the source \\citep{2009ATel.2002....1S} at a barycentred period of 12.44~s indicating that this object hosts an accreting X-ray pulsar. Here we report the detailed analysis of the \\integ, \\swift, and \\xte\\ observations. We begin this paper by detailing the procedures employed for the data reduction. We then describe the results (refined position, X-ray spectral, and temporal analyses) in Sects. 3, 4, 5, and 6, and discuss them in Sect. 7. ",
        "conclusions": "We have reported here the results obtained from a multi-instrumental campaign dedicated to the X-ray properties of new \\integ\\ sources. The refinement of the X-ray position provided by the \\swift/XRT observations enabled us to identify a possible counterpart at infrared wavelengths. The differences of dereddened magnitudes (Table \\ref{tab:mag}) do not lead to any of the spectral types tabulated in \\citet{tokunaga00}. With \\nh=2.0$\\times 10^{22}$ cm$^{-2}$,  J$-$H=0.28 would indicate an F7 V or F8 I star. However the value of H$-$K$_{\\mathrm s}$ is inconsistent with both possibilities. With \\nh=2.73$\\times 10^{22}$ cm$^{-2}$, J$-$H seems too high for any spectral type, although we remark a marginal compatibility (at the edge of the errors on the magnitude) with an O9.5 V star \\citep[J$-$H=$-0.13$, H$-$K$_{\\mathrm s}$=$-$0.04][]{tokunaga00}.  This value is, however, a measure of the interstellar absorption through the whole Galaxy. Since \\igr\\ probably lies closer than the other end of the Galaxy, it is likely that the interstellar absorption along the line of sight is lower. In particular, we note that with \\nh=2.5$\\times 10^{22}$ cm$^{-2}$, we obtain J$-$H=$-0.04$ and H$-$K$_{\\mathrm s}$=0.05, which is very close to the values tabulated for a B3 I star \\citep[J$-$H=$-0.03$, H$-$K$_{\\mathrm s}$=0.03][]{tokunaga00}. In that case, the source would lie at a distance $d \\gtrsim 8$ kpc.\\\\ \\indent The X-ray behaviour of the source is indicative of an HMXB: the X-ray spectra are power law like in shape with a hard photon index. One of the spectra (Obs. S1 with the longest exposure) shows evidence for absorption in clear excess to the absorption on the line of sight, which may indicate some intrinsic absorption in this source. However, the evidence is marginal in the other spectra, which could suggest that the intrinsic absorption varies in this system, as has been observed in a number of HMXBs. A long-term periodicity was revealed in the \\swift/BAT data which confirms the binarity of the source \\citep{2009ATel.2008....1C}. We clearly detect an X-ray pulse at a period of 12.44~s, indicating the presence of an X-ray pulsar. We estimate a pulse fraction $P=40\\%$. Apart from millisecond X-ray pulsars, found in systems that are at the end of the evolutionary path of X-ray binaries (and are LMXBs), pulsars are young objects that are usually found in HMXBs. We also remark that the PSD of \\igr\\ shows a broadening at the bottom of the coherent pulsation.  Sidelobes and other noise features around coherent pulsation signals have been seen in other HMXBs \\citep[e.g.][in, respectively, \\object{4U 1626$-$67}, and \\object{XTE J0111.2$-$7317}]{kommers98,kaur07}. Such features can be produced by artificial effects (such as the finite length of the time intervals used to make the PSD), or they can be real if, e.g., a quasi-periodic oscillation (QPO) beats with the coherent signal \\citep{kommers98,kaur07}. In the case of \\igr, a possible QPO might be present at $\\sim0.035$~Hz (Fig. \\ref{fig:psd}), even if the quality of the data does not allow us to firmly establish its presence.  On the other hand, we cannot exclude that the $0.078$~Hz feature is itself a QPO at a frequency close to that of the coherent pulsation. Given the faintness of the source, we cannot conclude further on that matter.  \\\\ \\indent In fact, \\igr\\  lies in a region populated by Be-HMXBs in the so-called ``Corbet diagram'' \\citep{corbet86,2009ATel.2008....1C}, as demonstrated in Fig. \\ref{fig:corbet}. Note that this plot is the most recent update of the Corbet diagram for \\integ\\ sources as of June, 2009 (Bodaghee et al. 2009 in prep.). \\begin{figure} \\epsfig{file=pulse_orbit.ps,width=\\columnwidth} \\caption{Most recent version of the ``Corbet diagram'' including all sources detected with \\integ\\ (Bodaghee et al. 2009 in prep.). Squares represent Be-HMXBs, circles Supergiant-HMXBs (Sg-HMXB), circles in squares are SFXTs, filled symbols represent sources discovered by \\integ. The triangles are HMXBs of unknown nature (including \\igr). The positions of \\igr\\ and of the two SFXTs lying in the Be region of the plot are highlighted. } \\label{fig:corbet} \\end{figure} Recently, \\citet{arash07} further explored the parameter spaces occupied by high-energy sources. They noticed that Be-HMXB and supergiant HMXB also segregate in different parts of the \\nh\\ vs. orbital period and \\nh\\ vs. spin period diagrams.  The value of the absorption we obtained through our spectral analysis, combined with the values of the spin and orbital periods, make \\igr\\ lie in a region populated by Be-HMXBs in these diagrams as well. At first order, one can easily understand the segregation in these three different diagrams as results of the age, type of accretion, and probable type of orbit. Be systems are younger, and hence have eccentric orbits, longer orbital periods, and shorter spin periods, whereas supergiant systems are older, are mostly circularised with shorter orbital periods, and longer spin periods. In these latter systems, in addition, the compact object is embedded in the wind of the companion (the feed matter for accretion) which explains their (usually) higher intrinsic absorption.  In this respect, all our results point towards \\igr\\ being a Be-HMXB \\citep[the same conclusion, although only based on the Corbet diagram, was presented by][]{2009ATel.2008....1C} .\\\\ \\indent \\integ\\ has unveiled a new (sub-)type of supergiant HMXB, characterised by short ($\\sim$hours) and intense flares seen at X-ray energies: the so-called SFXTs. Models of SFXTs involve stochastic accretion of clumps from the (heterogeneous) wind of the supergiant \\citep[e.g.][]{zand05, negueruela06}, on top of longer, but fainter, periods of activity \\citep[e.g.][]{sidoli07}. In this respect, the detection of periods of short and intense flares in \\igr\\ with \\integ\\ (Fig. \\ref{fig:igrlite}), a behaviour typical of SFXTs, raises the possibility that this source belongs to this new class. Its position in the Corbet and the \\nh\\ vs. spin or orbital period diagrams (Fig.~\\ref{fig:corbet}) is not a definitive contradiction to this hypothesis, since systematic X-ray studies of SFXTs have shown that, contrary to other HMXBs, they seem to populate any part of the diagrams. In particular, \\object{IGR J18483$-$0311} and \\object{IGR J11215$-$5952}, the only two SFXTs for which both orbital and pulse periods are known, lie in the region of Be-systems in these representations (Fig. \\ref{fig:corbet}). The former source has a pulse period of 21.05~s, an orbital period of 18.5~d, and has a B0.5 Ia companion \\citep[see][and references therein]{rahoui08_18483}. The latter has a pulse period of about 187~s, an orbital period of about 165~d \\citep[e.g.][]{sidoli07, ducci09}, and is associated with a B supergiant \\citep{negueruela05}. To reconcile the fact that they are long period systems with a rather low absorption, an eccentric orbit is invoked. \\igr\\ could be the third member of SFXTs (with a known pulse period) having a long and eccentric orbit. In this case, this may point towards the existence of an evolutionary link between Be-HMXBs and eccentric SFXTs (Liu et al. 2009 in prep.). The quality of the data is not good enough to permit us to make any firm conclusions concerning the nature of the system, and only an identification of the optical counterpart's spectral type will resolve this issue."
    },
    "0910/0910.2616_arXiv.txt": {
        "abstract": "{} {In this paper we study the effects of a toroidal magnetic flux tube emerging into a magnetized corona, with an emphasis on large-scale eruptions.  The orientation of the fields is such that the two flux systems are almost antiparallel when they meet.}{We follow the dynamic evolution of the system by solving the 3D MHD equations using a Lagrangian remap scheme.}{Multiple eruptions are found to occur.  The physics of the trigger mechanisms are discussed and related to well-known eruption models.} {} ",
        "introduction": "The trigger mechanism for large-scale eruptions, such as coronal mass ejections (CMEs), is one of the main concerns of theoretical solar physics.  There are many different approaches to modelling solar eruptions, ranging from 2D analytical models to 3D numerical simulations.  A review of such models can be found in \\citet{forbes00} and further references are given in \\citet{dmac09a}.  One highly influential model for solar eruptions is the \\emph{breakout model} \\citep{antiochos99}.  This theory uses, as an initial condition, a multipolar configuration, which consists of four distinct flux systems separated by a null point.  This configuration is stressed by an \\emph{imposed} shearing at the lower boundary (normally taken to be the photosphere).  Reconnection across the current sheet at a stressed null changes the field geometry and weakens the tension of overlying field lines, converting them into nonrestraining field lines of neighbouring flux systems.  The continuation of this process can eventually lead to expulsion of a flux rope.  For more details of the breakout model, the reader is directed to \\citet{antiochos99}, \\citet{devore08} and \\citet{lynch08}.  Recently, \\citet{devore08} and \\citet{lynch08} have simulated magnetic breakout in 3D.  The first of these studies finds homologous confined eruptions.  The second study follows the topological evolution of a fast breakout CME.  \\citet{soenen09} simulate homologous CMEs in the solar wind in an axisymmetric 2.5D configuration.  The breakout studies mentioned above have been very useful in investigating the physical mechanisms of solar eruptions.  One point which they all have in common, however, is that the initial equilibrium is stressed by artificially imposed shearing motions.  Other models are also used to study eruptions in the solar atmosphere.  For example, \\citet{mackay06} model the large-scale coronal field as a series of non-linear force force-free equilibria.  Flux ropes form but not all of them settle into equilibrium.  Those that diverge from equilibrium are ejected.  For studying dynamic evolution in the atmosphere, another model is \\emph{flux emergence}.  This considers the early evolution of active regions, which, of course, are the sources of large-scale solar eruptions.  The standard practice in dynamic flux emergence experiments is to have a stratified solar atmosphere including the top of the solar interior.  A flux tube is placed in the solar interior and the system is left to evolve by itself, with no imposed flows.  A review of such models can be found in \\citet{archontis08rev}.  Some studies include a magnetized corona in their model.  \\citet{archontis05} and \\citet{galsgaard07} study the effects of a cylindrical flux tube emerging into a horizontal coronal field.  They consider different orientations for the magnetic fields and discuss the reconnection and high-speed jets that occur.  \\citet{maclean09} consider the same experiment but investigate the magnetic topology of the system.  In relation to solar eruptions, \\citet{manchester04} report on the eruption of a flux rope that forms during the emergence of a cylindrical flux tube into an non-magnetized corona.  \\citet{archontis08} study the emergence of a cylindrical flux tube into a magnetized corona and find that with a favourable orientation, a CME-like eruption is possible.  In the present paper, we shall consider a similar setup to \\citet{archontis08} but use a different model for the magnetic field.  This is a toroidal loop, placed in the solar interior, rather than a cylindrical one \\citep{hood09}.  A comparison of these two models is discussed in \\citet{dmac09b}.  This paper will show that multiple large-scale eruptions are possible from the same emerging region and will discuss the physics of the trigger mechanisms.  The outline of the paper is as follows: $\\S 2$ will describe the model setup and the initial conditions.  In $\\S 3$ we shall discuss the results of the eruption experiment and link the processes involved with eruption models.   $\\S 4$ will summarize the results. ",
        "conclusions": "In this paper we have demonstrated that multiple CME-like eruptions are possible from a toroidal flux tube emerging into a magnetized corona.  This combines and builds on the work of \\citet{dmac09b} and \\citet{archontis08}.  For the present study we consider a corona with a field that is almost antiparallel to the field of the emerging tube.  External reconnection at the apex of the emerging arcade weakens the tension of the coronal field.  With the expansion of the arcade, this reconnection becomes faster through time. Shearing, which occurs as part of the emergence process, induces internal reconnection in the arcade and produces a flux rope.  The continued emergence in combination with removal of the overlying coronal field eventually results in the expulsion of the flux rope. The mechanism for this eruption is similar to that of the breakout model.  One important difference, however, is that the external reconnection in this model does not take place at a null point. After the first eruption, continued emergence and, therefore, shearing results in the formation of a second flux rope.  This also erupts but the trigger mechanism cannot be directly linked to the breakout model, as with the first eruption.  Due to the reconnection of the first eruption with the corona, a weakened coronal field exists above the second rope when it forms.  Possible trigger mechanisms, such as runaway reconnection and the torus instability, have been suggested.  However, it is possible that the trigger for the second eruption is a combination of such mechanisms."
    },
    "0910/0910.1615_arXiv.txt": {
        "abstract": "In star formation, magnetic fields act as a cosmic angular momentum extractor that increases mass accretion rates onto protostars and in the process, creates spectacular outflows. However, recently it has been argued that this magnetic brake is so strong that early protostellar disks -- the cradles of planet formation -- cannot form. Our three-dimensional numerical simulations of the early stages of collapse ($\\lesssim 10^5$ yr) of overdense star--forming clouds form early outflows and have magnetically regulated and rotationally dominated disks (inside 10 AU) with high accretion rates, despite the slip of the field through the mostly neutral gas. We find that in three dimensions, magnetic fields suppress gravitationally driven instabilities which would otherwise prevent young, well ordered disks from forming. Our simulations  have surprising consequences for the early formation of disks, their density and temperature structure, the mechanism and structure of early outflows, the flash heating of dust grains through ambipolar diffusion, and the origin of planets and binary stars. ",
        "introduction": "Over the past decade numerical simulations have enabled the exploration of the central physical questions regarding the nature of star formation -- from the collapse of an initial dense gaseous molecular cloud core to the formation of a star and its associated protostellar disk and outflow, and the emergence of  a star's planetary system.  The formation of disks and jets during star formation is central to many of these issues, but very little is known about the earliest phases of their evolution.  How is the initial excessive angular momentum associated with the star's natal core removed  -- through magnetic braking \\citep{1994ApJ...432..720B} and then outflows \\citep{2006ApJ...641..949B,2007prpl.conf..277P}, or by spiral density waves in disks \\citep[e.g.~][]{2009arXiv0901.4325L}?  Do multiple stars form through gravitational fragmentation of cores or massive disks \\citep{HW2004b,2007MNRAS.377...77P,2008A&A...477...25H}?   What is the significance of outflows and jets that are launched before most of the mass has collapsed into the disk \\citep{2003MNRAS.341.1360L, 2006ApJ...641..949B}  in a wide variety of young stellar systems -- from brown dwarfs \\citep{2005Natur.435..652W, 2009ApJ...699L.157M} to massive stars \\citep{2007prpl.conf..245A, 2007ApJ...660..479B}?  Extraction of angular momentum by magnetic fields that thread the collapsing gas may be too efficient, according to recent two-dimensional  \\citep{2008ApJ...681.1356M} and three-dimensional axisymmetric \\citep{2008A&A...477....9H} simulations, preventing the formation of rotationally dominated disks, even when the effects of imperfect coupling of the field with the gas are included \\citep{2009ApJ...698..922M}. These results seemingly contradict the observations which clearly show that disks are present around most if not all young stars, even in environments in which the magnetic field is expected to be strong \\citep{1995AJ....109.1846H}. \\citet{2008A&A...477....9H} performed three-dimensional ideal MHD simulations of collapsing cores using a barotropic equation of state, concluding that no rotationally dominant structure is formed from 10 to 100 AU for highly magnetized cores.   \\citet{2008ApJ...681.1356M} performed two-dimensional ideal MHD simulations on collapsing singular isothermal toroids using a barotropic equation of state and an inner boundary at 6.7~AU  (effectively a sink particle), finding that even moderately magnetized disks could not form.   Such two-dimensional  models impose a high degree of mathematical symmetry  and therefore miss the formation of bars and spiral waves in disks.  We improve on previous results through this three-dimensional adaptive mesh refinement (AMR) investigation which includes a full treatment of the cooling \\citep{2006MNRAS.373.1091B} and the finite coupling of the magnetic field to the pre--stellar gas \\citep[ambipolar diffusion,][]{2008MNRAS.391.1659D}.  We find that ordered disk--like structures can emerge on scales $\\lesssim 10$ AU at early times ($\\lesssim 10^5$ yr) in magnetized systems. We have omitted a sink particle in this study as they are expected to affect the solution to within a few sink radii.  Without sinks simulations are indeed limited to early disks, although they offer a full solution to the region within 10 AU where the heated core forms. ",
        "conclusions": "In the early stages of a purely hydrodynamic collapse of a moderately rotating core, material joins the protostar by accreting through a chaotic series of bars and spiral waves. Our results show that in magnetized collapses however, the magnetic field suppresses these wave modes, and a small, regular disk appears at the earliest times. The weakening of magnetic control by ambipolar diffusion is insufficient to guarantee the formation of binary stars in typical cores with moderate rotation. Our results suggest that modest turbulent amplitudes ($>$ 10\\%) appear to be required. Regarding planet formation, the central density structure of early disks (Figure \\ref{fig:2}(e) falls off as $\\Sigma \\propto r^{-(1.7-2.5)}$ much more quickly with disk radius than do protoplanetary disk models at later times, wherein $\\Sigma \\propto r^{-1.5}$ or $r^{-1}$. As the collapse winds up the field early outflows appear -- even for partially coupled disks --  and feed angular momentum and mechanical energy back into the star forming neighborhood. The density and magnetic structure of ideal and partially coupled disks are quite similar -- the main difference is in the strength of the wound up magnetic field (at least an order of magnitude weaker in the ambipolar diffusions case).  Finally, we note that evidence for early localized intense drift heating near the disk at the accretion shock, which heats materials up to 1,000K and beyond in a localized region, may be preserved in the observed composition of comets and meteorites. The rapid heating, and subsequent cooling of those crystalline Mg--rich silicate materials that passed through the accretion shock and through the region of high ambipolar heating is reminiscent of heating events that must have occurred for some of these materials seen in cometary grains \\citep{2007prpl.conf..815W}.  These events have been attributed to shock heating by spiral waves out to 10 AU in disks \\citep{2005ASPC..341..849D}. As seen in Figure 1(b), only a portion of the accreted disk material will pass through this localized heated region. Subsequent radial turbulent mixing of this flash heated material with the bulk of material that passed through a more gentle heating environment could in principle contribute to the wide mixture of thermal histories preserved in cometary grain materials."
    },
    "0910/0910.3610_arXiv.txt": {
        "abstract": "We report observations of the remnant of Supernova 1987A with the High Resolution Camera (HRC) onboard the \\emph{Chandra X-ray Observatory}. A direct image from the HRC resolves the annular structure of the X-ray remnant, confirming the morphology previously inferred by deconvolution of lower resolution data from the Advanced CCD Imaging Spectrometer. Detailed spatial modeling shows that the a thin ring plus a thin shell gives statistically the best description of the overall remnant structure, and suggests an outer radius $0.96\\arcsec\\pm0.05\\arcsec \\pm0.03\\arcsec$ for the X-ray--emitting region, with the two uncertainties corresponding to the statistical and systematic errors, respectively. This is very similar to the radius determined by a similar modeling technique for the radio shell at a comparable epoch, in contrast to previous claims that the remnant is 10-15\\% smaller at X-rays than in the radio band. The HRC observations put a flux limit of 0.010\\,cts\\,s$^{-1}$ (99\\% confidence level, 0.08-10\\,keV range) on any compact source at the remnant center. Assuming the same foreground neutral hydrogen column density as towards the remnant, this allows us to rule out an unobscured neutron star with surface temperature $T^\\infty>2.5$\\,MK observed at infinity, a bright pulsar wind nebula or a magnetar. ",
        "introduction": "The core-collapse supernova \\object{(SN) 1987A} in the Large Magellanic Cloud was the brightest SN observed since the invention of modern telescopes, providing the best opportunity to study the last evolutionary stage of a massive star. Optical observations have revealed a triple-ring nebula centered on the explosion, consisting of an inner equatorial ring of radius 0.81\\arcsec-0.86\\arcsec and two larger outer rings \\citep{bkh+95,plc+95}. These are part of the hourglass-shaped circumstellar medium (CSM) ejected by the progenitor 20,000 years ago \\citep[see][]{mp07}. Since early 2004, the SN blast wave has begun to encounter the main body of inner ring \\citep{pzb+05,pzb+06}. This `big crash' has led to a drastic increase in soft X-ray emission that originates from the optically thin thermal plasma behind the shock. The X-ray spectrum is well-described by a two-component plane-parallel shock model in nonequilibrium ionization, with plasma temperatures $kT\\sim 2$ and 0.3\\,keV, corresponding to the fast and decelerated shocks, respectively \\citep{pzb+04,pzb+06}. As the blast wave propagates into the dense CSM, the soft X-ray emission traces the evolution of the forward shock, probing the CSM structure and its density profile.  The nearness of SN~1987A (51.4\\,kpc) allows us to resolve a supernova remnant (SNR) morphology at a very young age. In X-rays, this is only possible with the \\emph{Chandra X-ray Observatory}, the highest resolution X-ray telescope compared to any previous, current and even planned future X-ray missions. Previous \\emph{Chandra} observations of SNR~1987A were all carried out by the Advanced CCD Imaging Spectrometer (ACIS). However, the ACIS detector has a pixel size 0.492\\arcsec, comparable to the full width at half maximum (FWHM) of the mirror's on-axis point-spread function (PSF). Although the effective resolution can be slightly improved by the dithering of the spacecraft and by applying a sub-pixel imaging algorithm \\citep{tmm+01}, image deconvolution is still required to fully resolve the remnant structure \\citep[e.g.][]{bmh+00}. To determine whether artifacts are introduced by the complicated non-linear reconstruction process, the ACIS results need to be compared with higher resolution direct X-ray images. An analogy can be drawn from the radio imaging campaign of SNR~1987A. Since 1992, Australia Telescope Compact Array (ATCA) observations at 9\\,GHz have revealed detailed structure of the radio shell using the super-resolved technique, but it was not until the upgrade of the ATCA in 2003 that the higher resolution diffraction-limited images at 18\\,GHz provided a direct image of the radio morphology \\citep[see review by][]{gsm+07}.  The High Resolution Camera (HRC) onboard \\emph{Chandra}, consisting of two microchannel plate detectors, offers smaller electronic readout pixels (0.13175\\arcsec) than ACIS. These better sample the PSF, providing a more straightforward imaging process without the need for deconvolution. Despite a smaller effective area for HRC than ACIS and the lack of any spectral resolution, the SNR is now very bright in X-rays below 2\\,keV, at which the HRC has good sensitivity, making it an ideal instrument for morphological studies. In this Letter, we report a detailed analysis of the first HRC observation of SNR~1987A.  \\begin{figure*}[ht] \\centering \\begin{minipage}[c]{0.3\\textwidth} \\includegraphics[height=53mm, bb=5 49 544 503, clip=]{f1a.ps} \\end{minipage} \\hspace*{6.0mm} \\begin{minipage}[c]{0.3\\textwidth} \\includegraphics[height=53mm, bb=89 49 544 503, clip=]{f1b.ps} \\end{minipage} \\hspace*{-3.8mm} \\begin{minipage}[c]{0.3\\textwidth} \\includegraphics[height=53mm, bb=89 49 544 503, clip=]{f1c.ps} \\end{minipage} \\\\[-1pt] \\hspace*{0.1mm} \\begin{minipage}[c]{0.3\\textwidth} \\includegraphics[height=58.7mm, bb=5 1 544 504, clip=]{f1d.ps} \\end{minipage} \\hspace*{5.98mm} \\begin{minipage}[c]{0.3\\textwidth} \\includegraphics[height=58.7mm, bb=89 1 544 504, clip=]{f1e.ps} \\end{minipage} \\hspace*{-3.8mm} \\begin{minipage}[c]{0.3\\textwidth} \\includegraphics[height=58.7mm, bb=89 1 544 504, clip=]{f1f.ps} \\end{minipage} \\caption{(a) Raw HRC data of SNR~1987A taken on 2008 Apr 28-29. The image was binned into the HRC detector pixel and centered at RA=$\\mathrm{05^h35^m28^s}$, Dec=--69\\arcdeg16\\arcmin11.2\\arcsec\\ (J2000). (b) Raw ACIS data taken on 2008 Jan 9. The image is in 0.3-8\\,keV energy range and was binned into the ACIS detector pixel. (c) Super-resolved radio image at 9\\,GHz taken by ATCA on 2008 Apr 23 \\citep{ngs+08}. (d) Deconvolved HRC image using the dataset shown in panel (a). (e) Deconvolved ACIS image using the dataset shown in panel (b), in 0.3-8\\,keV energy range \\citep{rpz+09}. (f) HRC data in panel (a) smoothed to 0.4\\arcsec\\ to match the resolution of the radio image in panel (c). All panels are on the same spatial scale.\\label{f1}} \\end{figure*} ",
        "conclusions": "We have presented a detailed analysis of \\emph{Chandra} HRC observations of supernova remnant 1987A. The remnant has a similar size as the optical inner ring, and its morphology is statistically best-fitted by a thin ring plus a thin shell. This provides direct evidence that the dense inner ring of circumstellar medium has been overtaken the supernova blast wave. Our spatial modeling also indicates a similar size between the X-ray-- and radio-emitting regions, confirming the picture that the supernova forward and reverse shocks are closely located. The X-ray flux limit on any possible central source rejects a young neutron star, unless it has some fast cooling mechanisms or is obscured by a small accretion disk."
    },
    "0910/0910.1086_arXiv.txt": {
        "abstract": "{Abundance anomalies observed in globular cluster stars indicate pollution with material processed by hydrogen burning. Two main sources have been suggested: asymptotic giant branch (AGB) stars and massive stars rotating near the break-up limit (spin stars). We propose massive binaries as an alternative source of processed material. We compute the evolution of a 20\\Msun~star in a close binary considering the effects of non conservative mass and angular momentum transfer and of rotation and tidal interaction to demonstrate the principle. We find that this system sheds about 10\\Msun~of material, nearly the entire envelope of the primary star. The ejecta are enriched in  He, N, Na, and Al and depleted in C and O, similar to the abundance patterns observed in gobular cluster stars. However, Mg is not significantly depleted in the ejecta of this model.  In contrast to the fast, radiatively driven winds of massive stars, this material is typically ejected with low velocity. We expect that it remains inside the potential well of a globular cluster and becomes available for the formation or pollution of a second generation of stars. We estimate that the amount of processed low-velocity material ejected by massive binaries is greater than the contribution of AGB stars and spin stars combined, assuming that the majority of massive stars in a proto-globular cluster interact with a companion and return their envelope to the interstellar medium. If we take the possible contribution of intermediate mass stars in binaries into account and assume that the ejecta are diluted with an equal amount of unprocessed material, we find that this scenario can potentially provide enough material to form a second generation of low-mass stars, which is as numerous as the first generation of low-mass stars, without the need to make commonly adopted assumptions, such as preferential loss of the first generation of stars, external pollution of the cluster, or an anomalous initial mass function.  } ",
        "introduction": "For a long time star clusters have been considered as idealized single-age, chemically homogeneous stellar populations. However, it has recently become clear that many clusters show multiple main sequences and sub giant branches and extended horizontal branches \\citep[e.g.][]{Piotto+07}, implying the existence of multiple populations within one cluster.% In addition, large star-to-star abundance variations are found for light elements such as C, N, O, Na, and Al, while the composition of heavier elements (Fe-group and $\\alpha$-elements) seems to be constant.  Field stars with the same metallicity do not exhibit these abundance patterns \\citep[for a review see][]{Gratton+04}. These chemical variations have been interpreted as originating from the presence of both a ``normal'' stellar population, exhibiting abundances similar to field stars of the same metallicity and a second  population of stars formed out of material processed by hydrogen burning via the CNO-cycle and by the NeNa and MgAl chains \\citep[e.g.][]{Prantzos+07}.  {According to \\citet{Carretta+09},  50-70\\% of the stars in gloular clusters belong to the second population.}  Two sources of processed ejecta have been proposed: the slow winds of \\emph{massive AGB stars}, which enrich their convective envelopes with H-burning products \\citep{Ventura+01,Dantona+02,Denissenkov+Herwig03} and fast-rotating massive stars (we refer to these as \\emph{spin stars}), which are believed to expel processed material centrifugally when they reach break-up rotation \\citep{Prantzos+Charbonnel06,Decressin+07a}.% In this scenario a first generation of stars is formed out of pristine material. Their low-velocity ejecta are trapped inside the potential well of the cluster and provide the material for the formation of a second generation of stars. Although both proposed sources are promising, matching the observed abundance patterns and providing enough ejecta for the formation of a second generation that outnumbers the first generation have proven to be two major challenges. In this Letter we propose \\emph{massive binaries} as a candidate for the internal pollution of globular clusters.  \\begin{figure*} \\centering \\includegraphics[angle=-90, width=1.0\\textwidth, bb=50 50 283 616]{13205Fg1.eps} \\caption{ Composition of the slow ejecta of the modeled binary system (Sec.~3) as a function of the ejected amount of mass.  The mass fraction $X$ of the main stable isotope of each element is given relative to the initial mass fraction $X_{\\rm i}$, except for Mg where we added $^{24}$Mg$,^{25}$Mg and $^{26}$Mg.  The average $X_{\\rm av}$ and the most extreme mass fraction $X_{\\rm ex}$ are given in each panel on a logarithmic scale: $[X] \\equiv \\log_{10} (X / X_{\\rm i})$. For helium we show the absolute mass fraction $Y$ instead. Mass ejected during the first and second mass transfer phases is separated by a thin vertical line. \\label{fig:chem}} \\end{figure*} ",
        "conclusions": "We propose massive binaries as a source for the internal pollution of globular clusters. The majority of massive stars are expected to be members of interacting binary systems. These return most of the envelope of their primary star to the interstellar medium during non conservative mass transfer.  We show that there may be more polluted material ejected by binaries than by the two previously suggested sources: massive AGB stars and the slow winds of fast-rotating massive stars. After dilution with pristine material, as lithium observations suggest, binaries could return enough material to form a chemically enriched second generation that is as numerous as the first generation of low-mass stars, without the need to assume a highly anomalous IMF, external pollution of the cluster or a significant loss of stars from the unenriched first generation.  In addition to providing a new source of slowly-ejected enriched material, binary interaction also affects the previously proposed scenarios.  Binary mass transfer naturally produces a large number of fast-rotating massive stars that may enrich their surroundings even more.  Binary interaction will also affect the yields of intermediate-mass stars. Premature ejection of the envelope in 4-9\\Msun~stars will result in ejecta with less pronounced anti-correlations, as suggested in the AGB scenario. On the other hand, we expect that binary-induced mass loss may also prevent the dredge-up of helium-burning products.  For a detailed comparison of the chemical predictions of this scenario, binary models {for a range of masses and orbital periods} are needed and population synthesis models are essential to fullly evaluate the mass budget of the different sources.  {Finally, some peculiar feature, such as the apparent presence of distinct, chemically homogeneous subpopulations in  $\\omega$~Cen and NGC~2808 \\citep[e.g.][]{Renzini08} deserves further attention. }"
    },
    "0910/0910.0175_arXiv.txt": {
        "abstract": "The derivation of nebular abundances in galaxies using strong line methods is simple and quick. Various indices have been designed and calibrated for this purpose, and they are widely used. However, abundances derived with such methods may be significantly biased, if the objects under study have different structural properties (hardness of the ionizing radiation field, morphology of the nebulae) than those used to calibrate the methods.  Special caution is required when comparing the metallicities of different samples, like, for example, blue compact galaxies and other emission line dwarf galaxies, or samples  at different redshifts. ",
        "introduction": "Why talk about nebular abundances in a symposium devoted to stellar populations in galaxies? The ultimate goal of stellar populations studies is to understand the evolution of galaxies.  In principle, the determination of  the metallicities of the stellar populations gives an evolutionary view of the chemical enrichment of galaxies. Such studies are just beginning (see e.g. Vale Asari et al. 2009 and contribution to this symposium).  For the moment, however, the overwhelming majority of studies on the chemical composition of galaxies  rely on global stellar metallicity indices (for early-type galaxies),  and on nebular abundances (for star-forming galaxies) derived from emission-line studies. Note that the abundances derived by one or the other method do not have the same meaning: Stellar indices provide an integrated (and weighted) metallicity of the different generations of stars while the chemical composition of nebulae is the result of mixing with the interstellar gas of the metals ejected by \\emph{all} the present and past stellar generations and is the same as that of the presently forming stars.  In this communication, we would like to draw attention to widely overlooked biases in the determination of nebular abundances. We will restrict to the so-called ``strong line methods'' for abundance determination (see e.g. Stasi\\'nska 2004 for a general introduction on nebular abundance determinations). After presenting a number of abundance calibrators,  We will show how their use may, in certain cases,  lead to erroneous statements. ",
        "conclusions": "Strong line methods to derive nebular abundances are very easy to apply but they are prone to systematic errors. In principle, they should be used only for objects  whose HII regions have the  same structural proprties as those of the calibrating samples. This recommendation is not easy to follow, but at least one should be aware that using the same calibration for different samples may produces important biases.  In particular,  claims on differences in oxygen abundance \\begin{itemize} \\item between samples of galaxies with different chemical evolution histories \\item between samples of galaxies with different star formation histories \\item  between samples of galaxies at different redshifts (observational selection effects may play a role) \\end{itemize} should be taken with, at a minimum, a grain of salt."
    },
    "0910/0910.0033_arXiv.txt": {
        "abstract": "Deep \\textit{Swift} UV/Optical Telescope (UVOT) imaging of the Chandra Deep Field South is used to measure galaxy number counts in three near ultraviolet (NUV) filters (uvw2: 1928 \\AA, uvm2: 2246 \\AA, uvw1: 2600 \\AA) and the $u$ band (3645 \\AA).  UVOT observations cover the break in the slope of the NUV number counts with greater precision than the number counts by the Hubble Space Telescope (HST) Space Telescope Imaging Spectrograph (STIS) and the \\textit{Galaxy Evolution Explorer} (\\textit{GALEX}), spanning a range from $21 \\lesssim m_{AB} \\lesssim 25$.  Number counts models confirm earlier investigations in favoring models with an evolving galaxy luminosity function. ",
        "introduction": "Galaxy number counts as a function of magnitude provide direct constraints on galaxy evolution in both luminosity and number density. Number counts in the UV, in particular, can help trace the star formation history of the universe.    Until recently obtaining faint galaxy number counts in the UV has been difficult due to the small areas surveyed \\markcite{Gardner00,Deharveng94,Iglesias04,Sasseen02,Teplitz06}({Gardner}, {Brown}, \\&  {Ferguson} 2000; {Deharveng} {et~al.} 1994; {Iglesias-P{\\'a}ramo} {et~al.} 2004; {Sasseen} {et~al.} 2002; {Teplitz} {et~al.} 2006).  While the {\\sl Galaxy Evolution Explorer (GALEX)} has allowed for the measurement of UV galaxy number counts over a wide field of view \\markcite{Xu05}({Xu} {et~al.} 2005, $\\sim 20$ deg$^2$), the confusion limit of {\\sl GALEX} restricts the magnitude range covered to 14 to 23.8 $m_{AB}$.  The deepest UV number counts from {\\sl HST\\ } range from $m_{AB}$= 23 to 29 over an extremely small field of view of $\\sim 1.3$ square arcminutes.  Here, we present galaxy number counts obtained in 3 near UV filters (1928 \\AA, 2246 \\AA, 2600\\AA) as well as in the $u$ band (3645\\AA) obtained using the {\\sl Swift\\/} UV/Optical Telescope \\markcite{UVOT}(UVOT; {Roming} {et~al.} 2005).   Deep exposures were taken of a 289 square arcminute field of view overlapping the Chandra Deep Field South \\markcite{CDFS}(CDF-S; {Giacconi} {et~al.} 2002) allowing for the measurement of number counts from $m_{AB}$ = 21 to 26.  UVOT data covers the break in the slope of the NUV number counts with greater precision than the existing  {\\sl GALEX} and {\\sl HST} number counts.  We use the UVOT number counts to explore the evolution of star-forming galaxies out to $z\\sim1$. ",
        "conclusions": ""
    },
    "0910/0910.2200_arXiv.txt": {
        "abstract": " ",
        "introduction": "Using simple hydrogen atoms, stars forge the entire range of atomic species we observe in Nature. Towards the end of their lives, stars eject these newly forged atoms into interstellar space. Out of this enriched interstellar gas, new stars, their planets and even life form. Stellar mass loss is the motor that makes stars and galaxies change over cosmic time. Yet, when and how mass is lost from stars is still an open debate. Stars more than eight times more massive than the Sun end their lives in spectacular explosions known as supernovae. Stars less than eight times more massive than the Sun lose up to 90\\% of their mass when they become giants, shortly after running of out hydrogen in their cores. The mass that is lost eventually forms beautiful nebulae around the mother star, objects that we call planetary nebulae; a misnomer acquired because they looked like the small circles of planets to 17th century astronomers. Modern telescopes have revealed they exhibit complex shapes (butterfly, multiple lobes emerging from disks, jets and bullets, objects pretty enough to fill coffee table books!). These shapes are in need of an explanation: how do stars lose quite so much mass and why is this mass not residing in a more or less spherical distribution around the star?  Simple round shapes have been understood in terms of the Interacting Stellar Winds (ISW) model \\citep{kwok82}. In this model, planetary nebulae (PNe) are formed when stellar material surrounding the central star (CSPN), originating in a slow wind from the asymptotic giant branch (AGB) progenitor, is swept up into a nebular shell by a hot, fast, ionizing post-AGB wind from the white dwarf CSPN. The great majority of PNe are not spherical and many theories have been introduced to explain the wide variety of observed shapes. Many shapes can be reproduced in models by introducing an asymmetry into the slow wind \\citep{kahn85,balick87}. Later theories have concentrated on the origins of the asymmetry, considering stellar rotation and/or magnetic fields \\citep{garcia-segura99,frank04,garcia-segura05}. Most recently though, such effects have been shown to be difficult to maintain in single AGB stars, leaving binary central star systems as the primary progenitors of PNe (see, for example, \\cite{demarco09} for a full review and discussion of the binary progenitor hypothesis).  There are several cases though where only the outer shells show a departure from symmetry. This has been postulated to arise from an interaction with the interstellar medium (ISM). In this review, we will consider the current knowledge in this area and the previous work which has led up to the interpretation of the PN--ISM interaction as a four stage evolutionary process, including a period of 'rebrightening', typically occurring late in the lifetime of a PN. We will review this particular part of the process, including the dominating effect of the earlier AGB phase of evolution, and consider the immediate progeny of rebrightened PNe, objects which can bear little resemblance to their origins. The observational support for each stage will also be summarized here for the first time. Finally, we will discuss how rebrightening changes the lifetime estimates of typical PNe and how a sizable sample of interacting PNe may change estimates of the PN population and the understanding of ISM structure on a galactic scale. ",
        "conclusions": "Simulations have shown that interaction with the ISM strongly affects the death of a PN, even one with an average velocity central star. The structure of the AGB wind around the CSPN is changed from a wall into a bow shock during the preceding phase of evolution. This bow shock and the ram pressure stripping of material into the accompanying tail then dominates the evolution of the PN. Rebrightening and rejuvenation occur, on a short timescale for high speed CSPN, as the PN shell interacts with the bow shock. Following this, the PN is stripped downstream in turbulent regions moving down the tail, resulting in the central star leaving its PN.  Further study of rebrightened PN and their surroundings can provide considerable information regarding several other research avenues. As highlighted in the opening paragraph, tracing the ejected mass in circumstellar structures such as bow shocks, tails and late-evolution turbulent regions swept down the tail can reveal much about the way AGB and postAGB stars lose mass. Such information as the position of the bow shock and the parameters of the tail can also reveal, taking into account the method of formation, not only historical stellar mass-loss rates and local ISM conditions, but also details of galactic orbits for the first time. The few objects that have observed lengthy tails, including Mira, HFG1 and Sh 2-68, are currently uniquely placed to derive such orbits. Future surveys of cool dust around AGB stars should reveal many more tails allowing a deeper investigation of a meaningful sample.  Once we have a meaningful sample of rebrightened PNe, we should also be able to see how the observable lifetimes of PNe are extended. Current statistical population estimates, which depend on estimated lifetimes, indicate a larger population than we have currently detected, but by extending that lifetime, population estimates accordingly reduce and we are likely to see much more of an agreement between the numbers we actually observe, which is believed to be fairly complete within the Milky Way, and the predicted population.  The model of PN--ISM interaction then can account for many characteristics of PNe and explain highly complex, disrupted objects such as Sh 2-68, but it does not reproduce several observed details. In particular, the fragmentation of the nebular rim is not reproduced by the simulations. More advanced models must be able to reproduce this, at least in the cases of high speed CSPN as it seems common at this end of the velocity distribution, e.g. Sh 2-188. No theoretical investigation has yet been performed employing 3D MHD simulations which include the galactic magnetic field. This would seem a natural next step and as Dgani and Soker point out, this may rapidly provide an explanation of particularly fragmentary shells. In a greater challenge to the model, the existence of ancient, effectively spherical PNe with very few signs of ISM interaction are difficult to explain within the model context, especially if they have displaced central stars which rules out an alignment between viewing angle and direction of motion. The model of PN--ISM interaction has come along way and revealed the phenomenon of rebrightened PNe, but there are a still a number of open questions to be solved and no doubt future observations will reveal many more."
    },
    "0910/0910.3206_arXiv.txt": {
        "abstract": "We present, for the first time, a clear $N$-body realization of the {\\it strong mass segregation} solution for the stellar distribution around a massive black hole. We compare our $N$-body results with those obtained by solving  the orbit-averaged Fokker-Planck (FP) equation in energy space. The $N$-body segregation is slightly stronger than in the FP solution, but both confirm the {\\it robustness} of the regime of strong segregation when the number fraction of heavy stars is a (realistically) small fraction of the total population. In view of recent observations revealing a dearth of giant stars in the sub-parsec region of the Milky Way, we show that the time scales associated with cusp re-growth are not longer than $(0.1-0.25) \\times T_{rlx}(r_h)$. These time scales are shorter than a Hubble time for black holes masses $\\mbul \\lesssim 4 \\times 10^6 M_\\odot$ and we conclude that quasi-steady, mass segregated, stellar cusps may be common around MBHs in this mass range. Since EMRI rates scale as $\\mbul^{-\\alpha}$, with $\\alpha \\in [\\frac{1}{4},1]$, a good fraction of these events should originate from strongly segregated stellar cusps. ",
        "introduction": "The distribution of stars around a massive black hole (henceforth MBH) is a classical problem in stellar dynamics \\citep{1976ApJ...209..214B,1977ApJ...211..244L}. The observational demonstration of the existence of nuclear stellar clusters (henceforth NSCs)---as revealed by a clear upturn in central surface brigthness---in the centers of galaxies makes its study ever more timely. A number of NSCs in coexistence with a central MBH have recently been detected \\citep{2009MNRAS.397.2148G} suggesting that NSCs around MBHs, like the one in the center of the Milky Way, may be quite common.  The renewed interest in this theoretical problem is thus motivated by the observational data in NSCs and, in particular, the very rich and detailed data available for the stars orbiting the Galactic MBH. At the same time,  the prospects for detection of  gravitational waves (GWs) from extreme mass ratio inspirals (henceforth EMRIs) with future GW detectors such as the {\\it Laser Interferometer Space Antenna} (LISA) also urge us to build a solid theoretical understanding of sub-parsec structure of galactic nuclei. In fact, EMRI rates will depend strongly on the stellar density of compact remnants as well as on the detailed physics within $O$($0.01$pc) of the hole, which is the region from which these inspiralling sources are expected to originate \\citep{2005ApJ...629..362H}.  \\cite{1976ApJ...209..214B} have shown, through a kinetic treatment that, in the case all stars are of the same mass, this quasi-steady distribution takes the form of power laws, $\\rho(r) \\sim r^{-\\gamma}$, in physical space and $f(E) \\sim E^p$ in energy space ($\\gamma=7/4$ and $p=\\gamma-3/2=1/4$). This is the so-called {\\it zero-flow solution} for which the net flux of stars in energy space is precisely zero. \\cite{2004ApJ...613L.109P} and \\cite{2004ApJ...613.1133B} were the first to report $N$-body realizations of this solution, thereby validating the assumptions inherent to the Fokker-Planck (FP) approximation---namely, that scattering is dominated by uncorrelated, $2$-body encounters  and, in particular, dense stellar cusps populated with stars of the {\\it same mass} are robust against ejection of stars from the cusp.  The properties of stellar systems that display a range of stellar masses are only very poorly reproduced by single mass models.  It is well known from stellar dynamical theory that when several masses are present there is mass segregation---a process by which the heavy stars accumulate near the center while the lighter ones float outward \\citep{1987degc.book.....S}. Accordingly, stars with different mass get distributed with different density profiles. By assuming a stellar population with two mass components, \\cite{1977ApJ...216..883B}---hencefort BW77---generalized their early cusp solution and argued heuristically for a scaling relation $p_L = m_L/m_H \\times p_H$ that depends on the star's mass ratio only. However, they obtained no general result on the inner slope of the heavy objects; nor did they discuss the dependence of the result on the component's number fractions. On the other hand, it was shown long ago by \\cite{1969A26A.....2..151H} that the presence of a mass spectrum leads to an increased rate of stellar ejections from the core of a globular cluster, but he did not include the presence of a MBH at the center. H\\'enon's work raises the question as to whether {\\it multi-mass} stellar cusps, obtained from the solution of the FP equation, are robust against ejection of stars from the cusp. Ejections---due to strong encounters---are {\\it a priori} excluded from the FP evolution, even though they could occur in a real nucleus. Furthermore, even if cusps were shown by $N$-body results to be robust against stellar ejections (and we show that they are in this Letter),  BW77 scaling  cannot be valid for arbitrary number fractions.  Recently \\cite{2009ApJ...697.1861A}---henceforth AH09---stressed this latter point and have shown via FP calculations that, indeed, in the limit where the number fraction of  heavy stars is realistically small, a new solution that they coined {\\it strong mass segregation} obtains with density scaling as $\\rho_H(r) \\sim r^{-\\alpha}$, where $\\alpha \\gtrsim 2$. They have shown that there are two branches for the solution parametrized by $\\Delta = \\frac{N_HM_H^2}{N_LM_L^2} \\cdot \\frac{4}{3+M_H/M_L}$. The $weak$ branch, for $\\Delta > 1$ corresponds to the scaling relations found by BW77; while the $strong$ branch, for $\\Delta < 1$, generalizes the BW77 solution. There is a straightforward physical interpretation. In the limit where heavy stars are very scarce, they barely interact with each other and instead sink to the center due to dynamical friction against  the sea of light stars. Therefore, a quasi-steady state forms in which the heavy star's current is not nearly zero and thus the BW77 solution does not hold. As $\\Delta$ increases, self-scattering of heavies becomes important and the resulting quasi-steady state forms with a nearly zero current for stars of all masses, so BW77 solution is recovered.  For all these reasons,  it is fundamental to verify the Bahcall-Wolf solution---as well as its Alexander-Hopman generalization---with $N$-body integrations. There has been a surprisingly small number of $N$-body studies of multi-mass systems around a MBH \\citep{2004ApJ...613.1143B,2006ApJ...649...91F}, and none of them reported the occurence of strong mass segregation.  In this Letter we use direct $N$-body integrations to show for the first time that:  (i) strong mass segregation is a robust outcome of the growth of stellar cusps around a MBH when $\\Delta < 1$; (ii) BW77 solution is recovered when $\\Delta > 1$; (iii)  as a corollary, we conclude that the rate of stellar ejections from the cusp is too low to destroy the high density cusps around MBHs---even though ejections from the cusp {\\it do} occur.  Furthermore, having validated the FP formalism, we proceed to use it to estimate the time scales for cusp re-growth starting from a wider range of models. \\cite{2009arXiv0909.1318M} obtained, for Milky Way type nucleus, times in large excess of a Hubble time. We use a FP formalism which, in contrast with that of the latter author, is tailored to follow the simultaneous evolution of the cusp of different stellar masses without any restrictions with respect to the values of $f(E)$ or $\\rho(r)$. With our FP solutions we show that, for $\\mbul \\lesssim 5 \\times 10^6 M_\\odot$, the times for re-growing  stellar cusps are shorter than a Hubble time. Our results clearly suggest that strongly segregated stellar cusps around MBHs in this mass range may be quite common in NSCs and should be taken into account when estimating  EMRI event rates. ",
        "conclusions": "Our results show that strong mass segregation is a {\\it robust} outcome from the growth of stellar cusps around MBHs. We have used $N$-body integrations with two masses---light and heavy components representing main sequence stars and stellar BHs respectively---, and compared the results with those obtained with the FP formalism. The broad agreement between both methods validates the FP description of  the {\\it bulk} properties of time-evolving stellar distribution around a MBH---and its underlying assumptions. The differences of quantitative detail are the subject of another work in preparation.  Using the FP equation to study cusp growth under a variety of initial conditions purported to represent cored nuclei, we have shown that the time scales associated with cusp re-growth are clearly shorter than a Hubble time for nuclei with MBHs in the mass range $\\mbul \\lesssim 5 \\times 10^6 M_\\odot$---even though the relaxation time, as estimated for a single mass stellar distribution, exceeds a Hubble time in the upper part of this mass range. Therefore, our work strongly suggests that quasi-steady---strongly segregated--- stellar cusps may be common around MBHs with masses in this range.  EMRIs of compact remnants will be detectable by LISA precisely for MBHs in this mass range \\citep{2006PhRvD..74b3001D,2007CQGra..24..113A,2007PhRvD..75b4005B}. Estimates for event and detection rates by LISA costumarily assume that the stellar cusps are in steady state \\citep{2006ApJ...645.1152H,2006ApJ...645L.133H}. But recent observations reveal a dearth of giants inside $1$ pc from $SgrA^*$ and raise the possibility that cored nuclei are common---this scenario has been thoroughly explored by \\cite{2009arXiv0909.1318M}.  Our results strongly suggest that stellar cusps can re-grow in less than a Hubble time. The existence of cored nuclei still remains plausible though---especially for nuclei with MBHs in the upper part of the mass range---, since time scales are still quite long ({\\it e.g.} $6$ Gyr in Milky Way type nuclei). However, since EMRI rates scale as $\\mbul^{-\\alpha}$, $\\alpha \\in [\\frac{1}{4},1]$,  and re-growth times are $\\lesssim 1$ Gyr for $\\mbul \\lesssim 1.2 \\times 10^6 M_\\odot$, we still expect that a substantial fraction of EMRI events will originate from segregated stellar cusps. Finally, indirect observations alone will reveal whether there is a ``hidden'' cusp of old stars and their dark remnants around $SgrA^*$ \\citep{2005ApJ...622..878W,2009ApJ...703.1743P}.  \\ \\ \\ \\ \\"
    },
    "0910/0910.3947_arXiv.txt": {
        "abstract": "X-ray observations of several active galactic nuclei show prominent iron K-shell fluorescence lines that are sculpted due to special and general relativistic effects. These observations are important because they probe the space-time geometry close to distant black holes. However, the intrinsic distribution of Fe line strengths in the cosmos has never been determined. This uncertainty has contributed to the controversy surrounding the relativistic interpretation of the emission feature. Now, by making use of the latest multi-wavelength data, we show theoretical predictions of the cosmic density of relativistic Fe lines as a function of their equivalent width and line flux. We are able to show unequivocally that the most common relativistic iron lines in the universe will be produced by neutral iron fluorescence in Seyfert galaxies and have equivalent widths $< 100$~eV. Thus, the handful of very intense lines that have been discovered are just the bright end of a distribution of line strengths. In addition to validating the current observations, the predicted distributions can be used for planning future surveys of relativistic Fe lines. Finally, the predicted sky density of equivalent widths indicate that the X-ray source in AGNs can not, on average, lie on the axis of the black hole. ",
        "introduction": "\\label{sect:intro} The X-ray spectrum of many active galactic nuclei (AGNs) exhibits a strong emission feature at an energy of 6.4~\\kev\\ that is due to the fluorescence of weakly ionized iron in a dense and relatively cold medium \\citep[e.g.,][]{nan89,pou90,np94}. The strength of the emission feature requires that the region responsible for the fluorescence subtend about half the sky as seen from the illuminating X-ray source \\citep[e.g.,][]{gf91}. As the variability properties of AGN place the X-ray source within 10~Schwarzschild radii from the black hole \\citep[e.g.,][]{gra92,utt07}, it is likely that relativistic effects may alter the observed line shape \\citep{fab89,laor91} allowing it to be a powerful probe of space-time curvature. Several examples of relativistically broadened \\fe\\ lines have been detected in the spectra of AGNs over the last decade \\citep[e.g.,][]{tan95,fab02,rn03,jmm07,fab09}. These detections are leading to measurements of black hole spins and accretion disk dynamics \\citep{br06,rf08}.  The strength of the Fe line emission depends on only four parameters: (i) the shape of the illuminating continuum \\citep{bfr02,br02}, which in AGNs is a power-law with photon index $\\Gamma$; (ii) the abundance of iron (relative to the solar value) in the accretion disk \\citep{bfr02}, $A_{\\mathrm{Fe}}$; (iii) the ionization parameter of the illuminated disk \\citep{bfr02}, $\\xi=4\\pi F_{X}/n_{\\mathrm{H}}$, where $F_{X}$ is the illuminating flux, and $n_{\\mathrm{H}}$ is the density of the illuminated slab; and (iv) the reflection fraction $R$, a measure of the relative strength of the reflection component in the observed spectrum (this is related to the covering factor of the accretion disk). Recently, multi-wavelength survey data has been able to measure correlations between each of the first three of these parameters to $\\lambda=L_{\\mathrm{bol}}/L_{\\mathrm{Edd}}$, the Eddington ratio of AGNs \\citep{rye09,ith07,nt07}. Here, $L_{\\mathrm{bol}}$ is the bolometric luminosity of an AGN, and $L_{\\mathrm{Edd}} = 1.3\\times 10^{38}\\ (M_{\\mathrm{BH}}/M_{\\odot})$~erg~s$^{-1}$ is the Eddington luminosity of a black hole with mass $M_{\\mathrm{BH}}$ and $M_{\\odot}$ is the mass of the Sun. In addition, various surveys have been able to measure the black hole mass distribution \\citep{net09} and its evolution with redshift \\citep{lab09}, as well as the luminosity function of AGNs as a function of redshift \\citep[e.g.,][]{ueda03}. We can then combine all this information to calculate the density and evolution of the \\fe\\ line from accretion disks. In the following we concentrate solely on the \\fe\\ line that arises from the inner accretion disk, and neglect the contribution from any narrow component \\citep[e.g.,][]{nan06}, or one that arises from reflection off a Compton-thick absorber \\citep[e.g.][]{my09}. The following $\\Lambda$-dominated cosmology is assumed \\citep{sper03}: $H_0=70$~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_{\\Lambda}=0.7$, and $\\Omega_{m}=0.3$. ",
        "conclusions": "\\label{concl} This work has shown, for the first time, the cosmic density and EW distribution of relativistic \\fe\\ lines. Sensitive observations of these line profiles are vital as they allow measurements of fundamental parameters such as the spin of the central black hole. Our results validate the previous and current results from \\textit{XMM-Newton} and \\textit{Suzaku} that have shown that the majority of relativistic \\fe\\ lines have EWs$< 100$~eV \\citep{nan07}. Now that models of the intrinsic line EW and flux distributions are available, detailed planning of future \\fe\\ surveys by \\textit{IXO} can be performed. The sensitivity of \\textit{IXO} will push spin measurements beyond the local universe and out to moderate redshifts, revolutionizing our understanding of the cosmic evolution of black holes.  Finally, we have shown that the extreme light-bending model predicts many more intense relativistic \\fe\\ lines than the $R=1$ model (Figs. 1 and 2), in disagreement with current observational constraints. This implies that, on average, the X-ray source in AGNs does not lie on the axis of the black hole where light-bending would be so extreme. It seems more likely that the X-ray source in average AGNs lies above the accretion disk at a radius of several $r_g$ from the black hole where light bending is less severe. Comparing data from future surveys of \\fe\\ lines with plots such as Figure~3 should allow a measurement of the average radial displacement of the X-ray source. Of course, this conclusion does not preclude the possibility that extreme light-bending occurs in a few rare sources."
    },
    "0910/0910.3481_arXiv.txt": {
        "abstract": "We present a fully consistent evolutionary disc model of the solar cylinder. The model is based on a sequence of stellar sub-populations described by the star formation history (SFR) and the dynamical heating law (given by the age-velocity dispersion relation AVR). The stellar sub-populations are in dynamical equilibrium and the gravitational potential is calculated self-consistently including the influence of the dark matter halo and the gas component. The combination of kinematic data from Hipparcos and the finite lifetimes of main sequence (MS) stars enables us to determine the detailed vertical disc structure independent of individual stellar ages and only weakly dependent on the IMF. The disc parameters are determined by applying a sophisticated best fit algorithm to the MS star velocity distribution functions in magnitude bins. We find that the AVR is well constrained by the local kinematics, whereas for the SFR the allowed range is larger. The model is consistent with the local kinematics of main sequence stars and fulfils the known constraints on scale heights, surface densities and mass ratios. A simple chemical enrichment model is included in order to fit the local metallicity distribution of G dwarfs. In our favoured model A the power law index of the AVR is 0.375 with a minimum and maximum velocity dispersion of 5.1\\,km/s and 25.0\\,km/s, respectively. The SFR shows a maximum 10 \\,Gyr ago and declines by a factor of four to the present day value of 1.5\\,$\\msun/\\mathrm{pc}^2/\\mathrm{Gyr}$. A best fit of the IMF leads to power-law indices of -1.46 below and -4.16 above 1.72$\\msun$ avoiding a kink at 1$\\msun$. An isothermal thick disc component with local density of $\\sim 6\\%$ of the stellar density is included. A thick disc containing more than 10\\% of local stellar mass is inconsistent with the local kinematics of K and M dwarfs. Neglecting the thick disc component results in slight variations of the thin disc properties, but has a negligible influence on the AVR and the normalised SFR. The model allows detailed predictions of the density, age, metallicity and velocity distribution functions of MS stars as a function of height above the mid-plane. The complexity of the model does not allow to rule out other star formation scenarios using the local data alone. The incorporation of multi-band star count and kinematic data of larger samples in the near future will improve the determination of the disc structure and evolution significantly. ",
        "introduction": "\\label{introd}  There is no 'concordance' Milky Way model available so far that describes the structure, kinematics, chemistry and evolution of the disc in great detail. The classical stellar density model of \\citet{bah80a,bah80b,bah84c} composed by a spheroid and an exponential disc with magnitude-dependent exponential scale heights is still widely used. \\citet{bah84a,bah84b} introduced a finite set of isothermal disc components and solved the Poisson and the Jeans equation for a dynamical equilibrium model of the vertical disc structure. At present the most sophisticated Milky Way model based on star counts is the so-called 'Besan\\c{c}on model' of the Galaxy developed and presented in a series of articles. In \\citet{rob03} a general description and the present status of the model is given. In this model the thin disc is composed by a sequence of stellar sub-populations with increasing age and velocity dispersion. The Besan\\c{c}on model has still some serious drawbacks in the construction of the thin disc concerning the density profiles of the components, the star formation history (SFR) and the initial mass function (IMF) that will be discussed below.  The main input functions to be determined in an evolutionary disc model are the SFR and the dynamical heating described by the age-velocity dispersion relation (AVR). The $\\mathrm{SFR}(t)$ of the Milky Way disc is still not very well determined. The main reason for that is the lack of good age estimators with corresponding unbiased stellar samples. Especially samples selected by colour cuts or by a magnitude limit are biased with respect to the age distribution, because there is an age-metallicity relation due to the chemical enrichment of the disc. Therefore not even the famous Geneva-Copenhagen sample of 14,000 F and G stars \\citep{nor04,hol09} (hereafter GCS1,GCS2) with well determined stellar properties and individual age estimates can be used to derive the star formation history directly by star counts. Systematic biases introduced by the method of GCS1 concerning age and stellar parameter determinations are discussed in \\citet{pon05} and \\citet{hay06}. The Besan\\c{c}on model of the Galaxy was developed and presented in a series of articles. In \\citet{hay97a,hay97b} the SFR of the thin disc is determined to be approximately constant using a series of isothermal components and applying an approximate method to achieve dynamical equilibrium. In further applications they used a constant SFR combined with a steep IMF in the sensitive mass regime of 1-3\\,$\\msun$ \\citep{rob03}. As a result, the mean age of the disc population, the scale heights and the surface densities of the stellar sub-populations are relatively small. \\citet{roc00} used chromospheric age determinations of late-type dwarfs. They apply stellar evolution and scale height corrections. The main result is the determination of enhanced star formation episodes over the lifetime of the disc. They exclude a constant SFR from chemical evolution models. In \\citet{fue04} the star formation history for open star clusters was determined. They found at least five episodes of enhanced star formation in the last 2\\,Gyr. Since star clusters contain only a small percentage of all disc stars, an extrapolation to the total (smoothed) SFR is not possible.  In recent years the Hipparcos stars with precise parallaxes and proper motions were used to determine the SFR with different methods. \\citet{her00} determined the SFR over the last 3\\,Gyr using isochrone ages. They found a series of star formation episodes on top of  an underlying smooth SFR. This result is complementary to our model, which gives the slow evolution of the smoothed SFR. \\citet{ver02} applied an inverse method to derive from the colour magnitude diagram the star formation history and age-metallicity relation. They used prescribed AVR and gravitational potential and fixed the IMF to be a power law over the whole stellar mass range. This ansatz results in a steep IMF and a rather young thin disc. Additionally they found strong maxima in the SFR at ages of $10^7$ and $2\\times10^9$\\,yr. In \\citet{bin00} and in \\citet{cig06} the scale height variation of main sequence stars was not taken into account. Therefore these models derive the local age distribution of stellar sub-populations with lifetimes larger than the age of the disc instead of the SFR. \\citet{bin00} determined the age of the solar neighbourhood to be $11.2\\pm 0.75$\\,Gyr, which includes the contribution of the thick disc. Both papers found an approximately constant local age distribution, which corresponds to a decreasing SFR according to the increasing scale height with increasing age of the stellar sub-populations. In a recent paper \\citet{aum09} derived the SFR using the updated Hipparcos data \\citep{vle07} and the kinematics of the GCS sample. They calculated in detail the number of main sequence (MS) stars as function of B-V colour based on a KTG93-like IMF and scale height corrections due to dynamical heating. Their favoured disc model has an exponentially decaying SFR with a decay timescale of 8.5\\,Gyr and an age of 12.5\\,Gyr. In the model the gravitational potential is fixed and the SFR depends strongly on the properties of the IMF by construction.  The presented paper (Paper I) starts a new attempt to develop a fully consistent disc model which is able to predict the properties of the Milky Way disc in great detail concerning density distributions, age distributions, kinematics and chemistry of all types of stars. The aim of this sequence of papers is to construct a smooth physically consistent disc model that can be extended to a global Galaxy model and allows detailed predictions of number densities and kinematics of stars of different types. The new model can be of some help for the construction of a 'concordance' model of the Milky Way. It can be used for the preparation of the future astrometric space mission Gaia that will measure with high precision positions, distances, proper motions and stellar parameters (temperature, surface gravity and chemical composition) of one billion stars \\citep{per01,bai05}.  We start with a local disc model of the solar cylinder, which is based on the SFR, the AVR and a chemical evolution function described by the age-metallicity-relation (AMR). The vertical density profiles and the corresponding scale height determinations of thin and thick disc stars are calculated in a self-consistent gravitational potential including the DM halo and the gas component. The model parameters are determined by a best fit procedure of the velocity distribution functions $f_\\mathrm{i}(|W|)$ of solar neighbourhood stars from B--K dwarfs. We use Hipparcos stars and at the faint end the Catalogue of Nearby Stars (CNS4), select the MS stars and divide these into a series of volume complete sub-samples in magnitude bins. The derived $f_\\mathrm{i}(|W|)$ of each sub-sample are  compared simultaneously to the model predictions. Details of the fitting procedure and the significance with respect to parameter variations are discussed in a future paper \\citep{jus09} (Paper II).  Since the velocity distribution function of MS samples is the average over the lifetime of these stars properly weighted by the SFR and the dilution due to the increasing scale height, the time resolution is strongly limited. Therefore we use only smooth input functions for the SFR, AVR and metal enrichment in the model. The resulting SFR, AVR and chemical evolution describe the long-term smoothed disc evolution. Similar disc models by fitting the vertical luminosity and colour profiles of edge-on galaxies instead of the kinematics have already been used successfully \\citep{jus96,jus06}.  The main advantages of this method is that it does not depend on individual stellar age determinations and that it is essentially independent of the shape of the IMF. A weakness is that we cannot confine the SFR strongly based on local data only. The application of the model to star counts of remote stellar populations will be of great help in this respect. In subsequent work the local model will be continuously extended and compared to large data samples taken from the Sloan Digital Sky Survey SDSS/SEGUE \\citep{aba09} and the Radial Velocity Experiment RAVE \\citep{zwi08} to further constrain the possible parameter range of the disc model.  A major restriction of a local model is that radial mixing of stars in the disc cannot be included easily. There are two main mixing processes discussed in the literature. The first one is directly connected to the dynamical heating process responsible for the AVR. The gravitational scattering process leads to a diffusion of orbits in velocity space (AVR) and in position (radial, tangential, vertical mixing). \\citet{wie96} discussed the radial diffusion of stellar orbits and the consequences for the AMR in some detail. The radial probability distribution function of the birth-places depends mainly on the age of the stellar sub-population and only weakly on the physical scattering process. For a 5\\,Gyr old population the typical radial width of the initial distribution is $\\pm 2$\\,kpc. The second mechanism is resonant scattering of circular orbits by spiral arms \\citep{sel02,ros08,sch09}. This mechanism changes the radial position of stars quickly, but the eccentricity of the orbits remain very small. The effect of resonant scattering may account for up to half the stars in the solar neighborhood  \\citep{ros08}, which shows the possible scale of the errors which may result if migration is ignored. Unfortunately the efficiency and the properties of resonant scattering in the Milky Way disc are still poorly known. In a future global disc model radial mixing may be included in a parametrized form. The local model presented in this paper is primarily a description of the 'status quo' of the local Milky Way disc quantified by the local age distribution (SFR), the kinematics (AVR) and the chemical composition (AMR). Therefore the relations between these function are not altered by radial mixing. But the physical interpretation of the SFR as a 'local star formation history' and the AVR as a 'local dynamical heating process' are weakened by radial mixing and require corrections. For the chemical evolution model the consequences are more severe. The tight connection of the enrichment, the SFR and the gas infall will be broken. The chemical enrichment of the gas decouples partly from the (apparent) enrichment of the stellar population. A more general local AMR including an intrinsic scatter would be the consequence as derived by \\citet{sch09}.  In section \\ref{model} the physics of the disc model is presented, section \\ref{observ} describes the observational data, in section \\ref{disc} the fitting procedure and the properties of the best fit models are given, section \\ref{summary} collects the main results and discussion. ",
        "conclusions": "\\label{summary}  We presented a new disc model for the thin disc in the solar cylinder based on a continuous  star formation history (SFR) and a continuous dynamical heating law (AVR) of the stellar sub-populations. This new model combines and improves the advantages of several different attempts to model the vertical structure in the solar neighbourhood: We used a sequence of isothermal sub-populations in dynamical equilibrium as \\citet{bah84c} and \\citet{aum09} did; We used the full velocity distribution functions $f|W|$ as \\citet{hol00} did, and not only the velocity dispersions (the AVR); We solved the Poisson equation self-consistently including the thick disc, gas and DM halo contribution as done in the Besan\\c{c}on model \\citep{rob03}. A chemical evolution model with reasonable gas infall rate, which was tuned to reproduce the local [Fe/H] distribution of G dwarfs, is included. This enables us to apply correct MS luminosities and lifetimes. Additionally our model is insensitive to the IMF, because it is based on the normalised velocity distribution functions of MS stars.  We determined pairs of (SFR, AVR) by a best fit of the local kinematics. The SFRs which are consistent with the MS velocity distribution functions show a decline factor below five down to unity (= const. SFR). The strongest feature that would distinguish between a constant and a declining SFR is a direct determination of the age distribution of low mass stars in the solar neighbourhood.  Despite the large variety of SFRs there is a strong correlation to the AVR. For each SFR the slope and maximum velocity dispersion of the AVR are well determined. For the AVR we find a power law with indices between 0.375 and 0.5. The range of models is consistent with the results of \\citet{bin00} for the local age distribution and \\citet{aum09} for the SFR.  Applying the stellar lifetimes and the new scale height corrections to the PDMF results in an IMF that shows only one break point at $1.7\\msun$ and a steep falloff at high masses.  The density profile of an isothermal thick disc component can be fitted very good by a sech$^\\alpha_\\mathrm{t}$ profile, where $\\alpha_\\mathrm{t}$ depends on the velocity dispersion. The most prominent effect of thick disc stars is the enhancement of the high velocity wings of K and M dwarfs in the range of 40--60\\,km/s. From that we can exclude a heavy thick disc with distinct kinematics and more than 10\\% contribution to the local stellar density. Changing the thick disc parameters leads to slight variations of the thin disc properties (mainly by assigning part of the stellar disc to the thick disc), but has a negligible influence on the normalised SFR and AVR of the thin disc.  A variety of predictions can be made from the new disc model. The density profiles of MS star sub-populations depend on the lifetime of the stars and are significantly different to density profiles of single age sub-populations. The shape is neither exponential nor sech$^2$ and can be characterised by the (half-)thickness and the exponential scale height. From the vertical density profiles MS stars number densities as function of colour and apparent magnitude can be predicted. Applying these number densities with observed 'Hess' diagrams from large surveys like the catalogues of the Sloan Digital Sky Survey SEGUE/SDSS enables us to restrict the parameters of the SFR further. We determined vertical gradients in the kinematics which will be tested with Radial Velocity Experiment (RAVE) data. Age and metallicity distributions of stellar sub-populations as a function of $z$ above the galactic plane are predicted.  The future plan is to extend the local model to a complete disc model of the Milky Way that provides a fully self-consistent connection of stellar densities and kinematics. Ultimately this kind of detailed model is essential to understand the large data sets as already available from SDSS and which are expected on a much higher level in amount and precision by PanSTARRS and the astrometric Gaia satellite mission."
    },
    "0910/0910.3162_arXiv.txt": {
        "abstract": "We have derived nebular abundances for 10 dwarf galaxies belonging to the M81 Group, including several galaxies which do not have abundances previously reported in the literature.  For each galaxy, multiple H~\\ii\\ regions were observed with GMOS-N at the Gemini Observatory in order to determine abundances of several elements (oxygen, nitrogen, sulfur, neon, and argon).  For seven galaxies, at least one H~\\ii\\ region had a detection of the temperature sensitive [OIII] $\\lambda$4363 line, allowing a ``direct\" determination of the oxygen abundance. No abundance gradients were detected in the targeted galaxies and the observed oxygen abundances are typically in agreement with the well known metallicity-luminosity relation.  However, three candidate ``tidal dwarf'' galaxies lie well off this relation, UGC~5336, Garland, and KDG~61.  The nature of these systems suggests that UGC 5336 and Garland are indeed recently formed systems, whereas KDG~61 is most likely a dwarf spheroidal galaxy which lies along the same line of sight as the M81 tidal debris field.  We propose that these H~\\ii~regions formed from previously enriched gas which was stripped from nearby massive galaxies (e.g., NGC~3077 and M81) during a recent tidal interaction. ",
        "introduction": "Due to their low dust content, simple kinematics, modest to negligible abundance gradients, and prevalence in the universe, dwarf galaxies are excellent objects to use as probes of galaxy evolution (e.g., Hunter \\& Gallagher 1985; Kobulnicky \\& Skillman 1997). Additionally, these small galaxies may be remnants of the building blocks that combined to form larger galaxies early in the history of the universe (e.g., Bullock \\& Johnston 2005).  Thus, understanding the chemical evolution of dwarf galaxies is critical for constraining galaxy evolution models.  Indeed, numerous studies have investigated dwarf galaxy abundances (e.g., Lequeux et al.\\ 1979;  Skillman et al.\\ 1989; Richer \\& McCall 1995; Kunth \\& \\\"Ostlin 2000 and references therein).  These studies have established the well known metallicity-luminosity relation, wherein more massive galaxies have higher oxygen abundances.  While this empirical trend has been measured repeatedly, its underlying cause(s) remains controversial.  On the one hand, greater retention of enriched gas by the deeper potential wells of more massive galaxies could produce the relation (e.g., Gibson \\& Matteucci 1997).  However, another possibility is that pristine gas is processed more efficiently by larger galaxies.  Thus, a combination of gas transport (i.e., inflows and outflows) and systematic variations in star formation efficiency could produce this trend (e.g., Dalcanton 2007).  At the same time, most galaxies do not exist in isolation.  Evolution and gas transport are influenced via galaxy-galaxy interactions; such influence is likely quite extreme in some cases.  If galaxies are hierarchically built structures, such interactions may be the {\\sl de facto} method of galaxy evolution.  Thus, interacting galaxies may further inform us as to how the composition of the universe is evolving.  Has interaction significantly altered the chemical composition of systems?  The nearby M81 Group (D $\\sim$ 4~Mpc) offers an ideal laboratory in which to investigate the signatures of chemical enrichment in a dynamic environment. This prominent group is one of the nearest galaxy associations to the Local Group.  Hence, even low mass group members may be observed and included in studies.  Additionally, the proximity of the M81 Group has permitted accurate distance determinations (Karachentsev et al.\\ 2002 and references therein) and resolved studies of the stellar populations (e.g., Weisz et al.\\ 2008). As recently as 300 Myrs ago, the primary galaxies of the M81 Group experienced a dramatic collision (e.g., van der Hulst 1979; Yun et al.\\ 1994).  The tidal HI debris still connects the three most massive galaxies involved in the event: M81, M82, and NGC 3077.  There have been detailed studies of this tidal debris in both HI and CO (Taylor et al. 2001, Chynoweth et al. 2008, Brinks et al. 2008).  Additionally, several clumps of new star formation, likely induced by the three-body interaction, have been identified (e.g., Davidge 2008).  Indeed, it has been suggested that some dwarf galaxies within the group have recently been formed through tidal interactions (Makarova 2002;  Sabbi et al. 2008; Weisz et al.\\ 2008).  Accordingly, one may hypothesize that the metallicity of such tidally formed galaxies could be elevated, compared to other galaxies of similar mass (e.g., Weilbacher et al. 2003); the gas from which they formed may be pre-enriched by the larger system from which the gas was stripped.  To investigate this possibilty, we have obtained spectra for 10 of the dwarf galaxies of the M81 Group, including some of the ``tidal dwarf\" candidates, to determine the intrinsic metallicity of these systems. Several of these galaxies do not have previously determined abundances in the literature.  We present these data and discuss their implications for the evolution of the M81 Group.  In \\S2 we present the observations and data processing.  The emission line measurements and abundance determinations are discussed in \\S3 and \\S4, respectively.  The results of the dwarf irregular galaxies are discussed in the context of the M81 Group in \\S5, while tidal dwarf galaxies are discussed in \\S6.  Lastly, \\S7 summarizes our findings.  We adopt 12$+$log(O/H) = 8.93 (Anders \\& Grevesse 1989) as the solar value of the oxygen abundance for the present discussion.   \\section {Observations} \\subsection{Optical Spectroscopy} Low-resolution spectra of ten M81 Group dwarfs were obtained\\footnote[1]{Program ID GN-2006A-Q-26.}  in queue mode with the Gemini Multiple Object Spectrograph (GMOS) instrument (Hook et al. 2004) on Gemini North during May 2006.  Galaxies were selected to span a large range of both star formation rate and optical luminosity.  Global parameters for the selected targets are listed in Table~\\ref{t:param}.  Masks were cut to place 5\\arcsec\\ long slitlets with a 1.5\\arcsec\\ slit-width on regions observed to emit significant flux in the H$\\alpha$ pre-imaging observations. Blue spectra were obtained using the  B600 (600 lines mm$^{-1}$) diffraction grating with a resolution of 0.45 $\\rm{\\AA}$ pixel$^{-1}$.  Red spectra were obtained using the R600 diffraction grating with a resolution of 0.47 $\\rm{\\AA}$ pixel$^{-1}$.   Spectral observations were typically 2$\\times$900s for both blue and red setups.  We used $4 \\times 2$ spectral-spatial binning, respectively, resulting in a spectral resolution of 1.8 $\\rm{\\AA}$ per binned pixel.  Observations were centered at 4500 $\\rm{\\AA}$ and 4550 $\\rm{\\AA}$ in the blue, and 7000 $\\rm{\\AA}$ and 7050 $\\rm{\\AA}$ in the red, to ensure full spectral coverage across the two detector gaps.  We did not use an order-blocking filter for any of the spectra.  This yielded complete spectral coverage from roughly 3500 $\\rm{\\AA}$ to 8000 $\\rm{\\AA}$ for the majority of slitlets.  The astrometric slit positions of emission sources presented in this work are listed in Table~\\ref{t:positions}.  Five apertures revealed regions of strong hydrogen emission yet displayed no sign of oxygen lines;  these sources were not analyzed as part of this work as they do not seem to be H~\\ii\\ regions, but their positions are also listed in Table~\\ref{t:positions}.  The position angle of the slitlets and the average parallactic angle for each pointing are also given in Table~\\ref{t:positions}.  Due to scheduling constraints, it was not always possible to line up the slits with the parallactic angle to reduce the light loss from atmospheric refraction.   The fields most affected by this misalignment are UGC 4459, UGC 5139, UGC 5918, and UGC 8201.  The spectra were reduced and analyzed using the IRAF\\footnote[2]{IRAF is distributed by the National Optical Astronomical Obsevatories.} GMOS package.  Reduction of spectra included bias subtraction and flat fielding based on observations of the quartz halogen lamp on the GCAL unit.  A spectral trace of a bright continuum source was used to define the trace for all slits in each field. The shape of this trace was consistent in all exposures of a given field.  Wavelength calibration was obtained by observations of CuAr arc lamps interspersed between observations of the target galaxies.  As all of the galaxies in this sample have relatively large angular extents, most filled a large fraction of the $5\\arcmin \\, \\times \\, 5\\arcmin$ GMOS-N field of view.  Thus, the edges of each 5\\arcsec\\ slit lie well within the glow of the diffuse ionized medium of the host galaxy.  Furthermore, some H {\\sc ii} regions are quite extended and completely fill these short 5\\arcsec\\ slitlets.  Therefore, we investigated whether standard local sky subtraction procedures could be used for these observations, given the concern that slits include both sky and diffuse ionized gas.  Indeed, local sky subtraction often over-subtracted strong lines relative to weaker emission lines, as expected since the diffuse ionized gas is at a different temperature than the compact H~\\ii\\ region.  Even in the most compact H {\\sc ii} regions observed, a 10\\% reduction in the flux of strong lines was found when the sky was measured locally.  Thus, a more accurate sky background was measured in slits located away from the main body of the target galaxy.  Average red and blue sky spectra were created by combining multiple one dimensional spectra from slits sampling the sky for every galaxy observed. While the same physical region was extracted from both the blue and red setups, the size of this extraction region varied from slit to slit.  Thus, the amount of sky background in each slit was different, as it depended on the spatial extent of the slit with significant line emission. To ensure the proper level of sky background was subtracted, spectral regions populated by sky lines and free of any substantial extraterrestrial emission line were used as reference regions (e.g., 7150--7300 \\AA).  Average sky spectra were scaled such that residuals in the reference regions were minimized when object and sky spectra were differenced.  Subsequently, the scaled backgrounds were subtracted from the one-dimensional H{\\sc ii} region spectra.  One-dimensional spectra were corrected for atmospheric extinction and flux calibration based on observations of the flux standard Feige 67 (Oke 1990). The flux standard was observed on each night science data were obtained.  A representative flux-calibrated spectrum is shown for each target galaxy in Figure \\ref{fig:spec}.  Note that red and blue continuum levels are in excellent agreement, despite observations being non-simultaneous.  This confirms stable sky conditions and indicates that the extraction regions were well matched in both the red and blue setups.  These example spectra are of similar quality to other spectra used for the abundance determinations of each target.  \\subsection{Line intensities} The strength of emission features were measured in the one-dimensional spectra and subsequently corrected for Balmer absorption and line of sight reddening.  Reddening estimates were determined by comparing the ratios of flux from Balmer emission lines.  Ratios of intrinsic line strength were interpolated from the tables of Hummer \\& Storey (1987) for case B recombination, assuming T$_e$=12500 K and N$_e$=100 cm$^{-3}$.  When the temperature sensitive [O {\\sc iii}] $\\lambda$4363 line was detected, temperatures derived from the [O {\\sc iii}] lines were used rather than 12500 K.  The Galactic reddening law of Seaton (1979) parameterized by Howarth (1983) was applied to derive the reddening function normalized at H$\\beta$, assuming R=A$_V$/E$_{B-V}$=3.1.  When the measured reddening coefficient, c$_{H\\beta}$, differed for the H$\\alpha$/H$\\beta$ and H$\\gamma$/H$\\beta$ line ratios, an underlying Balmer absorption of up to 2${\\rm \\AA}$ was applied.  The determined values of c$_{H\\beta}$ are reported in Table~\\ref{t:lineratiob}.  Reddening corrected line intensities measured from H {\\sc ii}  regions in the target fields are reported in Tables~\\ref{t:lineratiob} and \\ref{t:lineratior}.  The regions are grouped according to the field with which they are associated, as listed in the first column.    The error associated with each measurement is determined from measurements of the Poisson noise in the line measurement, error associated with the sensitivity function, Poisson noise in the continuum, read noise, sky noise, flat fielding calibration error, error in continuum placement, and error in the determination of the reddening.   Note that while the spectral coverage permitted measurements of lines redward of the [Ar {\\sc iii}] $\\lambda$7136 line, the numerous skylines made measurements in the red portion of the spectra very uncertain. Furthermore, both [O {\\sc i}] lines in our spectral region are coincident with strong atmospheric emission features.  Occasionally, the resulting profile of the sky-corrected [O {\\sc i}] $\\lambda$6300 and $\\lambda$6364 lines had broad wings indicative of over subtraction.  Thus, measurements of [O {\\sc i}] lines are not reported when sky subtraction did not yield trustworthy line profiles.  \\subsection{Diagnostic Diagrams} Since the red and blue spectra were not obtained simultaneously and, further, queue scheduling did not always permit observations along the parallactic angle, diagnostic diagrams were examined to verify we recovered reliable line ratios, and hence reliable abundances.  For example, the [N~\\ii]/H$\\alpha$ and [O~\\iii]/H$\\beta$ sequence is shown in Figure \\ref{fig:diag}a along with a theoretical curve from models of H~\\ii\\ regions (Baldwin, Phillips, \\& Terlevich 1981).  H~\\ii~regions from the current study are shown as circles and pentagons, while the triangles indicate H~\\ii~regions from a study of spiral galaxies (van Zee et al.\\ 1998).  Following McCall et al.\\ (1985), Figure \\ref{fig:diag}b shows the same diagnostic, replacing [O~\\iii]/H$\\beta$ with the sum of [O~\\ii]/H$\\beta$  and [O~\\iii]/H$\\beta$, commonly referred to as R$_{23}$.  Our data are in good agreement with both theoretical predictions and the general locus of galaxies in all diagnostic diagrams.  Intriguingly, H~\\ii~regions from the observed candidate tidal dwarfs, KDG 61, Garland, and UGC 5336, do not occupy the same parameter space as H~\\ii~regions from the dwarf irregulars in our sample.  Rather, both diagnostic trends indicate they resemble H~\\ii~regions observed in larger spiral galaxies.  We discuss these targets in more detail in \\S5.  \\section {Nebular Abundances} Determination of elemental abundances from spectra of ionized gas requires determination of (i) the electron density (n$_e$), (ii) the electron temperature (T$_e$), and (iii) an estimate of the abundance of atomic species in unobserved ionic states. In spectra where the [O {\\sc iii}] $\\lambda$4363 line is detected, we may directly follow the methodology of Osterbrock \\& Ferland (2006), since this line is highly sensitive to the electron temperature.  We adopt the standard practice of referring to this as the ``direct\" method (Dinerstein 1990).  When measurement of the weak line is impossible, we have employed a strong line method wherein an oxygen abundance is deduced from photoionization models and measurements of the strong [O {\\sc ii}] $\\lambda$3727 and  [O {\\sc iii}] $\\lambda$5007,4969 lines (e.g., Edmunds \\& Pagel 1984, McGaugh 1991).  Subsequently, a consistent electron temperature is deduced based on the strong-line oxygen abundance (van Zee et al.\\ 1997).  While the absolute measurement of this ``semi-empirical\" approach is less certain, it still provides robust relative abundances since relative abundances are less dependent on the electron temperature.  \\subsection{Direct Abundance Determinations} Osterbrock \\& Ferland (2006) carefully describe the procedure to determine the electron density and the electron temperature of an H~\\ii\\ region. We mention here only the most relevant points in the determination of T$_e$ and n$_e$.  The [S {\\sc ii}] $\\lambda$6717/6731 line ratio is sensitive to the electron density of the H {\\sc ii} region.  The majority of H {\\sc ii} regions studied here are within the low-density limit of this diagnostic ratio [$I(\\lambda6717)/I\\lambda6731)>1.35$].  Accordingly, an electron density of 100 cm$^{-3}$ was assumed for most H {\\sc ii} regions.  In the six cases where the line ratio did not lie below the low-density limit, electron densities were calculated from the observed [S {\\sc ii}] ratio using a version of the FIVEL program (De Robertis 1987).  Detection of the [O~\\iii] $\\lambda$4363 line in a spectrum permits a direct determination of the electron temperature of the ionized gas. We used the [O~\\iii] $\\lambda$4363 line for a temperature measurement only when the S/N in the line exceeded 2.8.  Following this threshold for detection, we were able to measure reliable reddening corrected [O~\\iii] $\\lambda$4363 fluxes for 26 H~\\ii~regions (56\\% of the sample). These 26 H~\\ii\\ regions were distributed among 7 of the 10 observed fields. From these fluxes, electron temperatures were determined for the O$^{++}$ region of the nebula. Subsequently, the electron temperature in the O$^+$ region was calculated from the O$^{++}$ electron temperature using the approximation of Pagel et al.\\ (1992).  With electron densities and temperatures in hand, emissivity coefficients for detected emission lines were determined using the FIVEL program (De Robertis et al. 1987).  While emission lines were detected for all dominant ionization states of oxygen, derivation of abundances for other elements requires accounting for the presence of atomic species whose ionization states were unobserved.  Nitrogen abundances were derived under the assumption N/O=N$^+$/O$^+$;  neon abundances were derived under the assumption Ne/O = Ne$^{++}$/O$^{++}$ (Peimbert \\& Costero 1969).  In the cases of both sulfur and argon, the analytical ionization correction factors of Thuan, Izotov, \\& Lipovetsky (1995) have been adopted.  When the [S~\\iii] $\\lambda$6312 line was not detected, the sulfur abundance was deemed to be too uncertain to report.  \\subsection{Semi-empirical Abundance Determinations} In the 20 H~\\ii\\ regions where the [O~\\iii] $\\lambda$4363 line could not be cleanly measured or was simply too weak to be detected, we are unable to determine the electron temperature via a direct approach.  However, the strong oxygen lines ([O~\\ii] $\\lambda$3729 and [O~\\iii] $\\lambda$4959,$\\lambda$5007) are sensitive to both oxygen abundance and electron temperature.  Therefore, a semi-empirical approach was used, where oxygen abundances were estimated from the strong line ratio (e.g., McGaugh 1991); we uniformly adopt an error of 0.20 for abundances determined via the semi-empirical method as discussed in van Zee et al.\\ (2006).  In this approach, where photoionization models are used as calibration, the geometry of the H~\\ii\\ region is represented by the average ionization parameter, \\=U, the ratio of ionizing photon density to particle density, as determined by the ratio of [O~\\iii] to [O~\\ii].  This additional parameter increases the spread of abundance estimates for a given R$_{23}$.  The theoretical photoionization models of McGaugh (1991) are shown in Figure \\ref{fig:mcg} along with the observed line ratios presented in this work; solid points indicate the [O~\\iii] $\\lambda$4363 line was detected while open points denote spectra where the line was unmeasurable.  In general, our data lie in the region predicted by theoretical models; however, three points fall off the model grid.  Two of these regions, KDG 61-9 and Garland-5, have detections of the [O {\\sc iii}] $\\lambda$4363 line; therefore, the electron temperature was determined directly from line ratios and does not rely upon photoionization models. The remaining point, UGC 5336-11, is from a supernova remnant.  As such, it would not be expected to follow predictions of pure photoionization but must be compared to models of shock ionization (see the discussion on UGC 5336 in \\S4).  The strong-line ratio of ([O~\\ii]+[O~\\iii])/H$\\beta$, or $R_{23}$, is a smooth function that depends on the electron temperature and oxygen abundance of a gas; however, $R_{23}$ is double valued.  Collisionally excited Lyman series emission accounts for the cooling of gas in a very low metallicity H~\\ii\\ region.  As the metal content increases, contributions to cooling from infrared fine structure lines become more important.  Around an oxygen abundance of one-third of the solar value (12$+$log(O/H) $\\sim$ 8.4), cooling becomes dominated by fine structure lines, causing the R$_{23}$ surface to fold over despite the increased metallicity.  The fold in the $R_{23}$ surface is shown in Figure \\ref{fig:mcg}.  The fold-degeneracy may be broken using the [N~\\ii] to [O~\\ii] line ratio as a diagnostic (e.g., Alloin et al.\\ 1979).  In the majority of cases here, H~\\ii\\ regions exhibited log([N~\\ii]/[O~\\ii])$<-1.0$.  Accordingly, the ``lower-branch'' and thus lower values for the oxygen abundance were adopted for those H~\\ii\\ regions.  The high abundance value was generally adopted for regions in the ``upper-branch'' with log([N~\\ii]/[O~\\ii])$>-1.0$.  However, some ($\\sim5$) of the observed H~\\ii\\ regions fell near this critical value.  In such instances, the two possible strong-line abundances were compared to abundances from other H~\\ii\\ regions within the field.  As the interstellar-medium chemical abundances of dwarf galaxies are generally spatially homogeneous (e.g., Kobulnicky \\& Skillman 1997; Lee \\& Skillman 2004; Lee et al.\\ 2005, 2006), the abundance branch yielding a more consistent value with other spectra was adopted; often these other spectra had detections of $\\lambda$4363, eliminating this uncertainty on their behalf.  The selected H~\\ii\\ regions in UGC 5692 were all faint targets with large errors in the [N~\\ii] to [O~\\ii] ratio.  Thus, this diagnostic line ratio did not present a clear resolution to the R$_{23}$ degeneracy. While no spectra from UGC 5692 yielded a solid [O~\\iii] $\\lambda$4363 detection, two of the apertures had weak detections of this line. These detections were more consistent with higher temperature H~\\ii\\ regions.  Furthermore, the temperatures derived from the [O~\\iii] line ratios yielded oxygen abundances consistent with the lower branch of the R$_{23}$ relation.  Therefore, the lower abundance was adopted for H~\\ii\\ regions in UGC 5692.  The strong-line abundance calibrations we have used are derived from zero-age H~\\ii\\ regions (McGaugh 1991).  As an H~\\ii\\ region ages, the characteristic spectrum and ionization parameter may evolve, introducing systematic abundance errors (e.g., Stasi\\'nska \\& Leitherer 1996; Olofsson 1997). van Zee et al.\\ (2006) showed that such corrections may be significant when log([O~\\iii]/[O~\\ii])$<-0.4$ in an H~\\ii\\ region.  Within the present sample, two regions in Garland (Garland-6 and Garland-9) and three regions in UGC 5336 (UGC5336-3, UGC5336-11, and UGC5336-12) lie within this highly uncertain parameter space.  Weisz et al. (2008) have shown concentrations of blue stars in both of these systems, indicative of star formation in the last 100 Myr.  Thus, these HII regions are associated with recent star formation, but we have no constraint on the character of the radiation field (i.e., reflective of a fully populated zero-age main sequence versus an evolved population of lower mass main sequence stars).  Lacking more information on the H~\\ii\\ regions sampled by Garland-6, Garland-9, and UGC5336-3, we have made no corrections for the age of the nebulae in the current analysis.  Abundance determinations from UGC5336-11 and UGC5336-12 are discussed further in \\S4.  \\subsection{Abundances} The abundance determinations for each H {\\sc ii} region are presented in Table~\\ref{t:abund}.  Abundances were calculated using the measured line strengths and corresponding emissivity coefficients as determined by FIVEL (De Robertis et al.\\ 1987).  Errors associated with the derived abundances were clearly dominated by the uncertainty of the electron temperature.  The abundances of nitrogen and the $\\alpha$-elements (Ne, S, and Ar) relative to oxygen are shown in Figure \\ref{fig:alpha}, along with solar values from Anders \\& Gervasse (1989).  At higher oxygen abundance, H~\\ii\\ regions do show an increased N/O ratio.  This relative increase is consistent with secondary production of nitrogen dominating the primary component at higher metallicities (see Figure 4 of Vila-Costas \\& Edmunds 1993).  As expected for primary elements, $\\alpha$/O trends appear to be constant as oxygen abundance increases.  The mean log (Ne/O) is $-0.82\\pm0.12$; the mean log (Ar/O) is $-2.21\\pm0.13$; the mean log (S/O) is $-1.50\\pm0.10$.  These mean values are in agreement with large samples of H {\\sc ii} regions in dwarf galaxies (e.g., Thuan et al.\\ 1995; Izotov et al.\\ 2004; Lee et al.\\ 2004; van Zee et al.\\ 2006).  By placing multiple apertures on targeted galaxies, we are able to investigate the uniformity of oxygen abundance across dwarf galaxies in the M81 Group.   Except for the candidate tidal dwarfs, the standard deviation of oxygen abundances within each galaxy is smaller than the error associated with the individual measurements.  This reaffirms previous results which have shown no detectable abundance gradients within dwarf galaxies (e.g., Kobulnicky \\& Skillman 1996, 1997).  While the suspected tidal dwarf galaxies Garland and UGC 5336 may show real spatial trends in oxygen enrichment, these galaxies are not representative of undisturbed dwarf galaxies.  These dwarf galaxies and their origins are discussed further in \\S6.1~--~\\S6.3.  The average oxygen abundance and log(N/O) for a given galaxy are tabulated in Table~\\ref{t:param}.  The average abundance was determined with a weighted average of the individual measurements. This relies upon the assumption that each H~\\ii\\ region is an independent  measure of the composition of the galaxy, which is taken to be uniform.  Abundances determined using the direct method are more certain.  Thus, when direct and semiempirical determinations were available, the strong line abundances were not used in the average.  It should be noted that all strong line abundances agree, within the errors, with the average abundances determined for their respective galaxies.  Also, due to their larger uncertainties, the semi-empirical abundance determinations would have very little effect in weighted averages.  \\section {Comparison with Previous Work} Previous studies of the M81 Group have reported abundances for half of the galaxies in this sample.  We compare our findings with those of previous studies below.  In general, we note that the new results presented here are consistent with previous observations of these systems.  {\\it UGC 4459.} -- Hunter \\& Gallagher (1985), Skillman, Kennicutt \\& Hodge (1989), Hunter \\& Hoffman (1999), Pustilnik et al.\\ (2003) and Saviane et al.\\ (2008) all obtained spectra of UGC 4459.  They determined an oxygen abundance, 12+log(O/H), of 8.59, 7.62, 7.79, 7.52$\\pm$0.08 and 7.83, respectively, where [O~\\iii] $\\lambda$4363 was only detected by Pustilnik et al.\\ (2003).  While the values of Skillman, Kennicutt \\& Hodge (1989) and Pustilnik et al.\\ (2003) are lower than the value determined here using the direct method (7.82$\\pm0.09$), they are in excellent agreement with the semi-empirical abundance we calculate for UGC 4459 (7.61$\\pm0.20$).  We adopt the mean oxygen  abundance determined from the direct method in this paper, which is in excellent agreement with both the Hunter \\& Hoffman (1999) and Saviane et al.\\ (2008) results.  {\\it UGC 5139.} -- A strong line abundance is reported for UGC 5139 based only on the detection of $\\lambda$4959 and $\\lambda$5007 (Miller \\& Hodge 1996).  To derive an abundance, they estimated the strength of [O~\\ii] based upon theoretical ratios of [O~\\iii]/[O~\\ii].  Our semi-empirical result for this same H~\\ii\\ region (7.75$\\pm0.20$) is in excellent agreement with their value of 12+log(O/H)$= 7.7\\pm0.3$.  Here, we adopt our direct-method average abundance of 8.00$\\pm0.10$.  {\\it UGC 4305.} -- Oxygen abundances of 12+log(O/H)=8.55, 7.92, and 7.71$\\pm$0.13 were reported by Hunter \\& Gallagher (1985), Masegosa et al.\\ (1991), and Lee et al.\\ (2003), respectively.  While the strong line abundance of Hunter \\& Gallagher (1985) appears to have been placed on the ``upper-branch\" of the strong line abundance models, the three HII regions where the [OIII] $\\lambda$4363 line was detected by Masegosa et al.\\ (1991) are in excellent agreement with our derived value of 7.92$\\pm$0.10.  We note the presence of strong He \\ii\\ $\\lambda$4686 emission in UGC 4305-11, which may indicate the presence of another ionizing source, such as a Wolf-Rayet star.  Such a composite spectrum could feature blends of lines which would not follow the general properties of a photoionized region.  Thus, we do not include UGC 4305-11 in our average abundance of UGC 4305.  {\\it UGC 5666.} -- Multiple H~\\ii\\ regions in UGC 5666 were studied by Masegosa et al.\\ (1991) and Miller \\& Hodge (1996), yielding 12 + log(O/H) = 8.09 and 8.06, respectively; these results are from a combination of direct and empirical methods.  While the current results are very tightly clumped together, with a resultant average of 7.93 $\\pm$ 0.05, they lie $\\sim0.1$~dex below the previously reported abundances.   UGC 5666 extends well beyond the $5\\arcmin \\, \\times \\, 5\\arcmin$ field of view of GMOS-N.  Therefore, diffuse galactic emission in the sky slit could have produced a systematic offset as discussed in \\S2.1.  However, it is also worth noting Masegosa et al.\\ (1991) obtained their data in 1984, before linear CCDs were widely employed, and thus a discrepancy of only 0.1 dex may not be significant.  {\\it UGC 5336.} -- An investigation of the optical counterpart of an X-ray source led to the acquisition of spectra from a U-shaped object north-east of UGC 5336 (Miller 1995).  Two of our apertures located in the UGC 5336 field lie upon this object.  While Miller (1995) reported anomalous H$\\alpha$/H$\\beta$ ratios for these regions and thus did not correct for line-of-sight reddening, our new observations yield normal H$\\alpha$/H$\\beta$ ratios indicative of foreground reddening.  In agreement with Miller (1995), we find the [S~\\ii] $\\lambda$6717,6731/H$\\alpha$ for slits UGC 5336-10 and UGC 5336-11 are 0.6 and 0.7 respectively, clearly above the criterion necessary for an object to be classified as a supernova remnant (Skillman 1985; Smith et al.\\ 1993).  Since UGC 5336-10 and UGC 5336-11 are not from purely photoionized H~\\ii\\ regions, proper abundance determinations require shock models.  Thus, diagnostic plots using [N~\\ii ]/H$\\alpha$, [S~\\ii] $\\lambda$6731/H$\\alpha$ and [O~\\iii]/H$\\beta$ have been constructed using the MAPPINGS III radiative shock ionization models (Allen et al.\\ 2008).  Under the assumption of a simple low density shock model, both diagnostics indicate 12+log(O/H)$\\sim8.4$.  Given that this value is consistent with the oxygen abundance results of photoionization models, we adopt the semi-empirical abundance determinations for these apertures.  The abundance determined from apertures sampling the supernova remnant (8.44$\\pm0.20$) as well as UGC 5336 (8.86$\\pm0.20$) are in excess of the findings of Miller (1995: 12+log(O/H)$\\sim$8.0).  While these two objects may not be physically associated, they appear to be linked by a common H I envelope.  Accordingly, we include measurements from all four apertures in the abundance determinations for UGC 5336. ",
        "conclusions": ""
    },
    "0910/0910.3212_arXiv.txt": {
        "abstract": "\\noindent We present a parametric analysis of the intracluster medium and gravitating mass distribution of a statistical sample of 20 galaxy clusters using the phenomenological cluster model of \\citeauthor{ascasibar08}. We describe an effective scheme for the estimation of errors on model parameters and derived quantities using bootstrap resampling. We find that the model provides a good description of the data in all cases and we quantify the mean fractional intrinsic scatter about the best-fit density and temperature profiles, finding this to have median values across the sample of 2 and 5 per cent, respectively. In addition, we demonstrate good agreement between \\rfiveh\\ determined directly from the model and that estimated from a core-excluded global spectrum. We compare cool core and non-cool core clusters in terms of the logarithmic slopes of their gas density and temperature profiles and the distribution of model parameters and conclude that the two categories are clearly separable. In particular, we confirm the effectiveness of the logarithmic gradient of the gas density profile measured at 0.04\\rfiveh\\ in differentiating between the two types of cluster. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.5328_arXiv.txt": {
        "abstract": "{The \\agile{} \\gray satellite accumulated data over two years on several blazars. Moreover, for all of the sources detected by AGILE, we exploited multi wavelength observations involving both space and ground based telescopes and consortia, obtaining in several cases broad-band spectral energy distributions (SEDs) which span from the radio wavelengths up to the TeV energy band.  I will review both published and yet unpublished \\agile{} results on \\gray blazars, discussing their time variability, their \\gray flare durations and the theoretical modeling of the SEDs. I will also highlight the GASP-WEBT and \\swi{} fundamental contributions to the simultaneous and long-term studies of \\gray blazars. ",
        "introduction": "Multi wavelength studies of \\gray active galactic nuclei (AGNs) date back to the late '70s and the early '80s with the COS--B detection of 3C~273 \\citep{Swanenburg78, Bignami81}. Nevertheless, the paucity of extragalactic \\gray source detected by SAS-2 and COS-B prevented systematic multi frequency studies. It was during the '90s, with the launch of \\cgro{}, that \\egret{} allowed to establish blazars as a class of \\gray emitters and to start multi wavelength studies of such sources. For a few sources, it was possible to study both the properties of the SEDs during different \\gray states, and the search for correlated variability at different bands, as for 3C~279 \\citep{Sed_Hartman2001, Var_Hartman2001}.  The recent launches of the \\agile{} and \\fermi{} satellites allowed the blazar community to observe a large fraction of the sky above 100~MeV, thanks to their wide ($\\sim 3$\\,sr) field of view (FoV), and to start a more effective multi wavelength approach in their spectral energy distribution investigation.  In the following, I will briefly introduce the \\agile{} satellite, and then I will focus on the \\agile{} results on the studies of \\gray blazars. Particular emphasis will be given to the importance of simultaneous (or at least, co-ordinated) multi wavelength observations, in order to study both the broad-band properties, and the correlations between the emission at different frequencies. ",
        "conclusions": "During its first year of sky monitoring, \\agile{} demonstrated the importance of its wide ($\\sim 3$\\,sr) field of view in detecting transient sources at high off-axis angles. Moreover, its unique combination of a \\gray detector with a hard X-ray monitor allowed us to study in detail the high energy portion of the blazar SEDs.  The synergy between the \\agile{} wide field of view, its fast response to external triggers, and the availability of a network of ground-based telescopes, allowed us to obtain a multi wavelength coverage for almost all the detected sources, and to investigate the physics of different classes of blazars.  Moreover, by means of long-term studies of selected objects, we were able to monitor both high and low \\gray states of different sources.  Finally, archival data analysis is in progress, and we start detecting dim and steady sources \\citep{Pittori2009AA:catal}."
    },
    "0910/0910.0482_arXiv.txt": {
        "abstract": "We describe a simple test of the spatial uniformity of an ensemble of discrete events.  Given an estimate for the point source luminosity function and an instrumental point spread function (PSF), a robust upper bound on the fractional point source contribution to a diffuse signal can be found.  We verify with Monte Carlo tests that the statistic has advantages over the two-point correlation function for this purpose, and derive analytic estimates of the statistic's mean and variance as a function of the point source contribution.  As a case study, we apply this statistic to recent gamma-ray data from the \\Fermi Large Area Telescope (LAT), and demonstrate that at energies above 10 GeV, the contribution of unresolved point sources to the diffuse emission is small in the region relevant for study of the WMAP Haze. ",
        "introduction": "Statistical tests of isotropy have a long history in astronomy.  A common question is ``What fraction of the observed emission could originate from unresolved point sources?''  For example, possible point source contributions to the extragalactic X-ray background were investigated by \\cite{1974MNRAS.166..329S}, and the small-angle power spectrum of the cosmic far infra-red background has been used to estimate the isotropic component \\citep{1996ApJ...473L...9K}.  More recently the Auger team tested the isotropy of ultra-high energy cosmic ray events by cross correlating with positions of known active galactic nuclei (AGN; \\cite{Cronin:2007zz,Abraham:2007si}) to provide information on their origin.  However in the absence of an appropriate external catalog, such cross-correlation methods cannot be used, motivating consideration of a more general approach.  In some cases a detector provides binned counts (e.g. pixels in a CCD); in other cases, photon event directions are reconstructed in some other way (e.g. a gamma-ray pair conversion telescope).  In the latter case, it is desirable to apply statistics that do not require binning of the data, as binning introduces additional arbitrary parameters into the problem.  In the limit of low flux density, where the mean density of photon events (hereafter, ``events'') is much less than one per PSF, explicit detection of point sources may become impractical, and estimation of the unresolved point source flux becomes especially difficult.  In some cases, the two-point correlation function, or some modified form \\citep[e.g.][]{Ave:2009id}, is used as a test.  However, the Fourier transform of a field of point sources has significant phase correlation, and a two-point function (or a power spectrum) discards this phase information.  Higher order correlation statistics capture it, but are somewhat cumbersome to use.  In the following, we describe a statistic that is easy to understand and evaluate, and is optimized to address this question, particularly in the case of fairly sparse data sets with (on average) $\\lesssim 1$ event per PSF circle.  The key insight is that if a substantial fraction of the photons come from point sources, it is much more likely that two photons appear within one PSF of each other than in the diffuse case.  This is true even if the expected integrated flux is of order one count.  Likewise, the number of PSF circles containing \\emph{no} counts is larger if point sources contribute.  These considerations motivate us to define a ratio between the fraction of ``isolated events'' and the fraction of ``empty circles.''  This ratio is very closely related to the fraction of diffuse emission, and can be calibrated with Monte Carlo simulations for specific choices of instrumental parameters and a putative luminosity function.  The two-point function, in contrast, is weighted by the density squared and is \\emph{not} proportional to the desired quantity.  In the following sections we define the statistic, estimate its variance, show how it behaves in various limits, generalize it to the case where many events appear in every PSF circle, and show a practical application to recently released data from the \\Fermi Gamma-ray Space Telescope.  \\begin{figure*} \\includegraphics[width=0.48\\textwidth]{plots/fsratio-diff100.ps} \\includegraphics[width=0.48\\textwidth]{plots/fsratio-diff85.ps} \\includegraphics[width=0.48\\textwidth]{plots/fsratio-diff50.ps} \\includegraphics[width=0.48\\textwidth]{plots/fsratio-diff10.ps} \\caption{\\label{fig:circles} In each panel, photon events (\\emph{dots}) are either isolated (\\emph{solid blue circles}) or not (\\emph{dashed blue circles}).  Random circles are either empty (\\emph{solid red}) or not(\\emph{dashed red}). In each case, 100 events and 100 random circles are shown, so the ratio, $R$, can be visualized here as the number of solid blue circles divided by the number of solid red circles.  In practise, one uses a large number of random circles to reduce noise.  The panels contain either no point sources (\\emph{upper left}), or 15\\% (\\emph{upper right}), 50\\% (\\emph{lower left}), or 90\\% (\\emph{lower right}) point source flux in sources located at (0.3,0.3) and (0.7,0.7). } \\end{figure*} ",
        "conclusions": "We have introduced a simple and easily calculable statistic that linearly traces the fraction of flux arising from diffuse emission, as opposed to unresolved point sources. The statistic is quite insensitive to even pronounced large-scale anisotropies in the diffuse emission, such as might originate from the proximity of a bright region or angular variation in the detector exposure. The linear response of this statistic to flux originating from point sources, and its smaller variance, make it superior to the two-point correlation function as a tracer of emission from unresolved point sources.  The sensitivity of the statistic to point source emission naturally depends on the luminosity function of the point sources, as a sufficiently steep power law extending to sufficiently small luminosities is strictly indistinguishable from diffuse emission. However, the statistic retains discriminatory power for spectral indices up to $\\alpha \\sim 3$, with a low-luminosity cutoff corresponding to an average of 0.1 counts, and assuming all point sources with average luminosity $\\gtrsim 10$ counts are resolved and removed. Known luminosity functions for astrophysical point sources generically have shallower slopes than this limit.  When the average number of events per PSF circle exceeds 1, the original form of the statistic breaks down: however, we have described a simple generalization suitable for this case, and demonstrated its efficacy. Increasing the number of counts by taking additional sky regions into account (i.e. without increasing the density of events) improves the variance by the usual $1/N_\\mathrm{event}$ Poisson factor. This statistic generalizes readily to higher dimensions; possible applications include the study of void statistics \\citep{1986ApJ...306..358F}.  As an example, we have applied this statistic to Class 3 (diffuse class) photon data from the \\Fermi LAT in the angular region relevant for study of the WMAP Haze, at energies of 10-100 GeV. We find that even with rather pessimistic assumptions for the point source luminosity function, at most $\\sim 15 \\%$ of the emission in this region can be attributed to unresolved point sources with average luminosities of $0.1+$ counts / year, and the results are consistent with $100 \\%$ diffuse emission.  We wish to acknowledge helpful conversations with Marc Davis, Josh Grindlay, Igor Moskalenko, Jim Peebles and Pat Slane. We thank the anonymous referee for helpful comments. TRS is supported by a Sir Keith Murdoch Fellowship from the American Australian Association.  \\begin{appendix}  \\begin{onecolumn}"
    },
    "0910/0910.3097_arXiv.txt": {
        "abstract": "{Wide-field \\textit{Spitzer} surveys allow identification of thousands of potentially high-$z$ submillimeter galaxies (SMGs) through their bright 24\\,$\\mu$m emission and their mid-IR colors.}{We want to determine the average properties of such $z\\sim$2 \\textit{Spitzer}-selected SMGs by combining millimeter, radio, and infrared photometry for a representative IR-flux ($\\lambda_{\\rm rest}\\sim 8\\,\\mu$m) limited sample of SMG candidates.}{A complete sample of 33 sources believed to be starbursts (``5.8\\,$\\mu$m-peakers'') was selected in the (0.5\\,deg$^2$) J1046+56 field  with selection criteria $F_{\\rm 24\\,\\mu m}$\\,\\textgreater\\,400\\,$\\mu$Jy, the presence of a redshifted stellar emission peak at 5.8\\,$\\mu$m, and $r^\\prime_{\\rm Vega}$\\,\\textgreater\\,23. The field, part of the SWIRE Lockman Hole field, benefits from very deep VLA/GMRT 20\\,cm, 50\\,cm, and 90\\,cm radio data (all 33 sources are detected at 50\\,cm), and deep 160\\,$\\mu$m and 70\\,$\\mu$m \\textit{Spitzer} data. The 33 sources, with photometric redshifts $\\sim1.5\\,-\\,2.5$, were observed at 1.2\\,mm with IRAM-30m/MAMBO to an rms $\\sim$0.7\\,-\\,0.8\\,mJy in most cases. Their millimeter, radio, 7-band \\textit{Spitzer}, and near-IR properties were jointly analyzed.}{ The entire sample of 33 sources has an average 1.2\\,mm flux density of $1.56 \\pm 0.22$\\,mJy and a median of 1.61\\,mJy, so the majority of the sources can be considered SMGs. Four sources have confirmed 4\\,$\\sigma$ detections, and nine were tentatively detected at the 3\\,$\\sigma$ level. Because of its 24\\,$\\mu$m selection, our sample shows systematically lower $F_{\\rm 1.2\\,mm}/F_{\\rm 24\\,\\mu m}$ flux ratios than classical SMGs, probably because of enhanced PAH emission. A median FIR SED was built by stacking images at the positions of 21 sources in the region of deepest \\textit{Spitzer} coverage. Its parameters are $T_{\\rm dust} = 37 \\pm 8$\\,K, $L_{\\rm FIR} = 2.5 \\times 10^{12}\\,L_{\\odot}$, and SFR\\,=\\,450\\,$M_{\\odot}$\\,yr$^{-1}$. The FIR-radio correlation provides another estimate of $L_{\\rm FIR}$ for each source, with an average value of  $4.1 \\times 10^{12}\\,L_{\\odot}$; however, this value may be overestimated because of some AGN contribution. Most of our targets are also luminous star-forming $BzK$ galaxies which constitute a significant fraction of weak SMGs at $1.7 \\lesssim z \\lesssim 2.3.$}{\\textit{Spitzer} 24\\,$\\mu$m-selected starbursts and AGN-dominated ULIRGs can be reliably distinguished using IRAC-24\\,$\\mu$m SEDs. Such ``5.8\\,$\\mu$m-peakers'' with $F_{\\rm 24\\mu m}$\\,\\textgreater\\,400\\,$\\mu$Jy have $L_{\\rm FIR}\\,\\gtrsim 10^{12}\\,L_{\\odot}$. They are thus $z \\sim 2$ ULIRGs, and the majority may be considered SMGs. However, they have systematically lower 1.2\\,mm/24\\,$\\mu$m flux density ratios than classical SMGs, warmer dust, comparable or lower IR/mm luminosities, and higher stellar masses. About 2000\\,$-$\\,3000 ``5.8\\,$\\mu$m-peakers'' may be easily identifiable within SWIRE catalogues over 49\\,deg$^2$.} ",
        "introduction": "Ultra-Luminous InfraRed Galaxies (ULIRGs, with $L_{\\rm FIR} \\gtrsim 10^{12}\\,L_{\\odot}$) are the most powerful class of star-forming galaxies. For 25 years, these prominent sources and their intense starbursts have been the target of many comprehensive studies, both locally \\citep[e.g., ][]{Sand96,Lons06,Veil09} and at high redshift \\citep[e.g., ][]{Blai04,Solo05}.  While local ULIRGs are relatively rare, submm/mm surveys with large bolometer arrays such as JCMT/SCUBA [James Clerk Maxwell Telescope/Submillimetre Common User Bolometer Array \\citep[]{Holl99}], APEX/LABOCA [Atacama Pathfinder Experiment/Large Apex Bolometer Camera \\citep[]{Siri09}] or IRAM/MAMBO [Institut de Radioastronomie Millim\\`etrique/Max-Planck Bolometer Array \\citep[]{Krey98}] have shown that the como\\-ving density of submillimetre galaxies (SMGs), which represent a significant class of high-redshift ($z\\sim1-4$) ULIRGs, is about a thousand times greater than that of ULIRGs in the local Universe \\citep[e.g.,][]{lefl05,Chap05}. They represent a major phase of star formation at early epochs and are also characterized by high stellar masses \\citep[e.g.,][]{Bory05}. They are thus ideal candidates to be the precursors of local massive elliptical galaxies \\citep[e.g.,][ hereafter Lo09, and references therein]{Blai02,Dye08,Lons08}.  Nearly all of the enormous UV energy produced by their massive young stars is absorbed by interstellar dust and re-emitted at far-infrared wavelengths, with their far-infrared luminosity ($L_{\\rm FIR}$) able to reach $10^{13}\\,L_{\\sun}$. However, despite the considerable efforts invested in mm/submm surveys, the total number of known SMGs remains limited to several hundred, and current observational capabilities are still somewhat marginal at many wavelengths. We thus still lack comprehensive studies of SMGs and their various subclasses at all wavelengths and redshifts and in various environments. Even their star formation rates (SFRs) remain uncertain in most cases because of a lack of direct observations at the FIR/submm wavelengths of their maximum emission. The identification of large samples of SMGs is important for investigating the properties of these galaxies (SFR, luminosity, spectral energy distribution [SED], stellar mass, AGN content, spatial structure, radio and X-ray parameters, clustering, etc.) on a statistical basis, as a function of their various subclasses, redshift, and environment. This is the main goal of the wide-field submm surveys planned with SCUBA2 and \\textit{Herschel}.  Although \\textit{Spitzer} generally lacks the sensitivity to detect SMGs in the far-IR, its very good sensitivity in the mid-IR allows the efficient detection of a significant fraction of SMGs in the very large area observed by its wide-field surveys, and in particular the $\\sim49\\,\\rm{deg}^2$ {\\it Spitzer} Wide-area Infrared Extragalactic (SWIRE) survey \\citep{Lons03}. From an analysis of a sample of $\\sim 100$ SMGs observed with \\textit{Spitzer}, Lo09 have estimated that SWIRE has detected more than 180 SMGs with $F_{\\rm 1.2\\,mm} > 2.5\\,$mJy per square degree at 24\\,$\\mu$m and in several IRAC bands from 3.6 to 8.0\\,$\\mu$m. However, the identification of SMGs among SWIRE sources is not straightforward, since it requires inferring FIR emission from mid-IR photometry in objects with various SEDs, especially as regards AGN versus starburst emission, and various redshifts.  We have therefore undertaken a systematic study of the 1.2\\,mm emission from the best SMG candidates among \\textit{Spitzer} bright 24\\,$\\mu$m sources, focusing on $z\\,\\sim$\\,2 starburst candidates.  In \\citet{Lons06} and Lo09 \\citep[see also][]{Weed06b,Farr08}, it is shown that selecting sources with a secondary maximum emission in one of the intermediate IRAC bands at 4.5 or 5.8\\,$\\mu$m provides an efficient discrimination against AGN power-law SEDs. In particular, 24\\,$\\mu$m bright ``5.8\\,$\\mu$m-peakers'' have a high probability of being dominated by a strong starburst at $z\\sim 2$, whose intense 7.7\\,$\\mu$m feature is redshifted into the 24\\,$\\mu$m band. A first 1.2\\,mm MAMBO study of a sample of $\\sim 60$ bright SWIRE sources (Lo09) has confirmed that such a selection yields a high detection rate at 1.2\\,mm and a significant average 1.2\\,mm flux density, showing that the majority of such sources are high-$z$ ULIRGs, probably at $z\\sim 2$. However, as described in Lo09, this sample was selected with the aim of trying to observe the ``5.8\\,$\\mu$m-peakers'' with the strongest mm flux over more than 10\\,deg$^{2}$. This was achieved by deriving photometric redshifts, estimating the expected 1.2\\,mm flux densities by fitting templates of various local starbursts and ULIRGs to the optical and infrared (3.6$-$24\\,$\\mu$m) bands, and selecting the candidates predicted to give the strongest mm emission. Therefore, the selection criteria of this sample were biased, especially toward the strongest 24\\,$\\mu$m sources and those in clean environments.  We report here the results of an analogous MAMBO study, but of a complete 24\\,$\\mu$m-flux limited sample of all SWIRE ``5.8\\,$\\mu$m-peakers'' in a 0.5\\,deg$^{2}$ region within the SWIRE Lockman Hole field, with $F_{\\rm 24\\,\\mu m} > 400\\,\\mu$Jy and $r^\\prime_{\\rm Vega} > 23$ (see Sec.\\ 2 for a precise definition of ``5.8\\,$\\mu$m-peakers'', which of course depends on the actual SWIRE data and limits of sensitivity and accuracy). This region was selected because of the richness in multi-wavelength data, in particular the exceptionally deep radio data at 20\\,cm \\citep[VLA,][]{OwMo09}, 50\\,cm (GMRT, Owen et al. in prep.), and 90\\,cm \\citep[VLA, ][]{Owen09}. Our study aims at characterizing the average multi-wavelength properties of these sources, their dominant emission processes (starburst or AGN), their stellar masses, and their star formation rates. We adopt a standard flat cosmology: $H_{0}$=71\\,km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_{M}$=0.27 and $\\Omega_{\\Lambda}$=0.73 \\citep{Sper03}. ",
        "conclusions": "The aim of this project was to determine the average properties of a complete 24\\,$\\mu$m flux limited sample of bright \\textit{Spitzer} sources selected to be starburst dominated at $z\\sim$2, using multi-wavelength data. The sample of 33 $z\\sim$2 SMG candidates was built with all the optically faint sources in a $\\sim$0.5\\,deg$^{2}$ area of the Lockman Hole SWIRE field meeting selection criteria based on MIPS/IRAC fluxes. These criteria are $F_{\\rm{24\\,\\mu m}}$\\,\\textgreater\\,400\\,mJy, a peak in the 5.8\\,$\\mu$m IRAC band due to redshifted 1.6\\,$\\mu$m stellar emission, and $r^{\\prime}_{\\rm{Vega}}$\\,$>$\\,23. The J1046+59 field was selected because of the availability of very deep radio observations at 20\\,cm and 90\\,cm with the VLA and at 50\\,cm with the GMRT. All sources in our sample are detected at 50\\,cm.  The entire sample has an average 1.2\\,mm flux density of 1.56$\\pm$0.22\\,mJy. However, the limited sensitivity allowed only four confirmed 4\\,$\\sigma$ detections, plus nine tentative 3\\,$\\sigma$ detections. Since the average 1.2\\,mm flux density, 1.56\\,mJy, corresponds to a 850\\,$\\mu$m flux density close to 4\\,mJy, about half of the sources may be considered SMGs. However, their redshifts range from $z\\sim$1.7$-$2.3, similarly to the sample in \\citet{Huan08}, but smaller than the redshift range covered by SMGs. The sample selected here is characterized by brighter 24\\,$\\mu$m flux densities, on average, than those of SMGs, and consequently shows systematically lower $F_{\\rm{1.2\\,mm}}$/$F_{\\rm{24\\,\\mu m}}$ ratios than classical SMGs. It is quite likely that our selection favours the SMGs with the brightest 24\\,$\\mu$m flux densities, due probably to enhanced PAH emission.    From stacking individual images of the sources, we are able to build the median FIR SED of our sample and estimate the corresponding $L_{\\rm{FIR}}$, SFR and $T_{\\rm{dust}}$ assuming a single temperature ``greybody'' model. The inferred values are $T_{\\rm{dust}}$\\,=\\,37$\\pm$8\\,K, $L_{\\rm{FIR}}$\\,=\\,2.5$\\times$10$^{12}$\\,$L_{\\odot}$, and SFR\\,=\\,450\\,$M_{\\odot}$\\,yr$^{-1}$. These estimates indicate that most of the sources are ULIRGs. However, estimates of $L_{\\rm{FIR}}$ for individual sources deduced from the IR-mm SED are highly uncertain due to the lack of flux measurements between 100\\,$\\mu$m and 500\\,$\\mu$m. The high quality radio data provide important complementary information on $L_{\\rm FIR}$ and the star formation rate, since the 1.2\\,mm/radio flux density ratio of the majority of individual sources is consistent with the FIR/radio correlation, which allows a derivation of $L_{\\rm{FIR}}$ and SFR from the radio flux, providing further confirmation that most of the selected sources are ULIRGs. The average value of  $L_{\\rm{FIR}}$  inferred from the FIR-radio correlation is 4.1$\\times$10$^{12}$\\,$L_{\\odot}$; however, this value may be overestimated because of an AGN contribution.   Stellar masses are estimated by modelling the optical-IR SED with stellar population synthesis models. They are of order a few 10$^{11}$\\,$M_{\\odot}$. Roughly scaling with the observed 5.8\\,$\\mu$m fluxes, they are similar to those of other samples of 24\\,$\\mu$m-bright, $z\\sim$2 \\textit{Spitzer} starbursts, and slightly higher than those of classical SMGs.  Overall, this sample appears similar to other samples of \\textit{Spitzer} $z\\sim$2 SMGs~\\citep[Lo09;][]{Youn09} in terms of millimetre emission, $L_{\\rm FIR}$, and SFR.      The complete radio detection of all sources provides a good estimate of the total star formation rate of such sources. They represent a significant fraction of all SMGs in the redshift range $z\\sim$\\,1.7-2.3 ($\\sim$10$-$15\\%). Most of these ``5.8\\,$\\mu$m-peakers'' are star-forming BzK galaxies with luminosities at the top of the luminosity distribution of sBzKs.  The surface density of ``5.8\\,$\\mu$m-peakers'' has been found to be  61\\,deg$^{-2}$ by \\citet{Farr06}. This is consistent with 33 sources in the 0.49\\,deg$^{2}$ of our field. We may thus estimate that 40\\,$-$\\,60 similar ``5.8$\\mu$m-peakers'' per square degree could be identified in the full SWIRE survey (49\\,deg$^{2}$). Most of them should be $z\\sim2$ starburst ULIRGs. At least half of them may be considered to be SMGs, including a small fraction of composite obscured AGN/starburst objects. Another significant fraction may be considered as SFRGs.  These results illustrate the power of deep multi-$\\lambda$ studies, especially with complete radio data, for analysing populations of powerful high-$z$ IR and submm sources. Such deep radio data are essential for disentangling starbursts and infrared-bright AGN, and for easily providing estimates of their star formation rates. We note especially the impressive complete detection of relatively weak $z\\sim$2 SMGs over 0.5\\,deg$^2$ in a single pointing of the GMRT at 610\\,MHz. As already proved by the analysis of SCUBA sources \\citep[e.g.,][]{Ivis02}, radio data are essential for identifying optical/near-IR counterparts and analysing submm surveys. This will be even more crucial for future surveys at the confusion limits of instruments like {\\it Herschel} at 300-500\\,$\\mu$m and SCUBA2 at 850\\,$\\mu$m. Even as we wait for EVLA and the new generation of SKA precursors, our results show that the GMRT at 610\\,MHz and even 325\\,MHz can already currently provide sensitivity well matched to wide {\\it Herschel} surveys.  It would be interesting to explore further whether the main properties which characterize this sample, i.e. strong MIR emission, radio activity, and high stellar mass, are related. Some of these properties are likely the result of biases introduced by our selection; however, this is unlikely to be the case for all of them, especially for the radio properties. In particular, a comparison between the starburst morphology (traced by young stars, dust, PAHs or CO emission, as measured by ALMA or {\\it JWST}) and the radio size would probe whether the radio emission is produced by the starburst or by an AGN, and whether the parameters of the starburst are different from those of most classical SMGs and reveal a different star formation regime.  We have several multi-wavelength observations planned or in progress for this sample to obtain better estimates of redshifts, dust temperatures, star formation rates, PAH luminosities, and AGN contributions, and thus constrain the dominant emission processes, and investigate the evolution and clustering properties of these sources."
    },
    "0910/0910.3742_arXiv.txt": {
        "abstract": "We present the results of three separate searches for \\HI\\ 21-cm absorption in a total of twelve damped Lyman-$\\alpha$ absorption systems (DLAs) and sub-DLAs over the redshift range $z_{\\rm abs}=0.86-3.37$. We find no absorption in the five systems for which we obtain reasonable sensitivities and add the results to those of other recent surveys in order to investigate factors which could have an effect on the detection rate: We provide evidence that the mix of spin temperature/covering factor ratios seen at low redshift may also exist at high redshift, with a correlation between the 21-cm line strength and the total neutral hydrogen column density, indicating a roughly constant spin temperature/covering factor ratio for all of the DLAs searched. Also, by considering the geometry of a flat expanding Universe together with the projected sizes of the background radio emission regions, we find, for the detections, that the 21-cm line strength is correlated with the size of the absorber. For the non-detections it is apparent that larger absorbers (covering factors) are required in order to exhibit 21-cm absorption, particularly if these DLAs do not arise in spiral galaxies. We also suggest that the recent $z_{\\rm abs} = 2.3$ detection towards TXS 0311+430 arises in a spiral galaxy, but on the basis of a large absorption cross-section and high metallicity, rather than a low spin temperature \\citep{ykep07}. ",
        "introduction": "Damped Lyman-$\\alpha$ absorption systems (DLAs) are believed to be the precursors of modern day galaxies, containing at least 80\\% of the neutral gas mass density of the Universe \\citep{phw05}. As the name suggests, DLAs are identified through their heavily damped absorption features, due to the large columns of neutral hydrogen ($N_{\\rm HI}\\geq2\\times10^{20}$ \\scm) through which the background quasar is viewed.  The 21-cm spin-flip transition is of interest since it traces the cool component of the gas, with the comparison of the 21-cm absorption strength to the total neutral hydrogen column giving the gas spin temperature ($T_{\\rm spin}$) for a fully absorbed emission region ($f=1$) [see Equ.~\\ref{enew}, Sect. \\ref{or}]. Searches for 21-cm absorption in DLAs exhibit a $\\approx50\\%$ detection rate, the detections occuring predominately at low redshift ($z_{\\rm abs}\\lapp1$), suggesting an increase in the $T_{\\rm spin}/f$ ratio with redshift, which could be due to higher spin temperatures and/or lower covering factors.  In order to shed light on which is the predominant factor, we have undertaken searches with the Green Bank (GBT) and Giant Metrewave Radio (GMRT) Telescopes. In particular, we aim to: \\begin{enumerate} \\item Test the hypothesis of \\citet{cmp+03} that 21-cm absorption should be readily detectable at high redshift, despite the large $T_{\\rm spin}/f$ ratios, by searching for 21-cm absorption in DLAs which lie towards compact radio sources. \\item Test the hypothesis of \\citet{ctp+07} that the spin temperature does not continue to increase with redshift, by searching in previously unsearched high redshift ($z_{\\rm abs}\\gapp3.2$) DLAs. \\item Test the line strength--metallicity correlation ($T_{\\rm spin}/f$\\,--\\,[M/H] anti-correlation) of \\citet{ctp+07}, which suggests that several, currently undetected, DLAs should be readily detectable in 21-cm absorption. \\end{enumerate} We have observed a dozen DLAs and sub-DLAs and in this paper we present and discuss the results in the context of the above issues. ",
        "conclusions": "\\subsection{Covering factors} \\subsubsection{Angular diameter distances and other effects}   \\label{copr}  As mentioned in Sect. \\ref{GBT06}, \\citet{cmp+03} suggested that, contrary to the high redshift equals high spin temperature argument of \\citet{kc02}, 21-cm absorption should be readily detectable at high redshift towards compact radio sources, where the covering factor is maximised. At the time of writing up the results presented in \\citet{cw06}, this was confirmed by the detection of 21-cm absorption in the highest redshift example to date, at $z_{\\rm abs}=2.347$ towards 0438--436 (\\citealt{kse+06}), which has a background radio source size of only $0.039''$ at 5 GHz \\citep{tml+98}. Following this, since submitting the proposal for the 2006 GBT observations, 21-cm absorption has been confirmed at yet higher redshift, $z_{\\rm abs}=3.386$ towards 0201+113 (\\citealt{kcl06}, see also \\citealt{dob96,bbw97}), as well as at $z_{\\rm abs}=2.289$ towards 0311+430 \\citep{ykep07}. Again, both of these occult compact radio sources of $<0.007''$ at 1.4 GHz \\citep{sbom90}\\footnote{\\citet{kcl06} report a de-convolved source size of $17.6\\times6.6$ mas$^2$ at 328 MHz.} and $1.36''\\times0.63''$, respectively. The latter was noted by \\citet{ykep07} to be unusual in its low spin temperature of $T_{\\rm spin}\\leq140$~K at such a high redshift, although this is not unusual in regard to the finding of \\citet{cmp+03,cw06}.  From Table \\ref{size} we see that five of our targets are towards radio source sizes of $\\theta_{\\rm QSO}\\lapp1''$ (recently found to be $0.07''$ for 0454+039 and $0.17''$ for 0528--250, \\citealt{klm+09}), although we have only meaningful $T_{\\rm spin}/f$ limits for three of these (0454+039, 0528--250 \\& 0758+475 -- Table~\\ref{res}). However, all of the GBT 2006 targets (Sect. \\ref{GBT06}) were selected before the arguments of \\citet{cw06} were formulated, and so all of the 2006 targets (and two of the three just mentioned), are all at high redshift, and will therefore be disadvantaged with respect to the DLA-to-QSO angular diameter distance ratios.  \\begin{figure*} \\centering \\includegraphics[angle=270,scale=0.73]{2-distance-z-hist-klm+09-actual.eps} \\caption{Top: The absorber/quasar angular diameter distance ratio versus the absorption redshift. Updated from \\citet{cw06}, where the symbols are as per Fig. \\ref{Toverf} and we include only the DLAs which have been searched to $T_{\\rm spin}/f\\geq100$~K (see main text). The iso-redshift curves show how $DA_{\\rm DLA}/DA_{\\rm QSO}$ varies with absorption redshift, where $DA_{\\rm QSO}$ is for a given QSO redshift, given by the terminating value of $z_{\\rm abs}$ (throughout this paper we use $H_{0}=71$~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_{\\rm matter}=0.27$ and $\\Omega_{\\Lambda}=0.73$). That is, we show the $DA_{\\rm DLA}/DA_{\\rm QSO}$ curves for $z_{\\rm em}$ = 0.5, 1, 2, 3 \\& 4. Bottom: The radio source size versus the absorption redshift. The symbols are colour coded according to the ratio of the frequency of the image to the redshifted 1420 MHz frequency: green -- $\\nu_{\\rm size}/ \\nu_{\\rm abs}\\leq3$, orange -- $3<\\nu_{\\rm size}/ \\nu_{\\rm abs}\\leq10$ and red -- $\\nu_{\\rm size}/ \\nu_{\\rm abs} > 10$. The hatched histogram represents the detections and the bold histogram the non-detections.} \\label{2-distance-z} \\end{figure*} Showing these in Fig. \\ref{2-distance-z} (top), we see that all but one of our target DLAs (for which we have limits), are indeed at angular diameter distance ratios of $DA_{\\rm DLA}/DA_{\\rm QSO}\\approx1$, due to their high redshift selection. For a given absorber and radio source size, this will disadvantage the effective covering factor (possibly giving the observed $\\overline{[T_{\\rm spin}/{f}]_{z \\geq 1}}\\approx2\\,\\overline{[T_{\\rm spin}/{f}]_{z < 1}}$, Sect. \\ref{oftpr}), although, as mentioned above, the three new high redshift detections occult compact radio sources, as shown in the bottom panel of Fig. \\ref{2-distance-z}, where we use the source sizes compiled in \\citet{cmp+03} updated with those given in \\citet{klm+09}\\footnote{\\label{foot7}Many of the source sizes are measured at (VLBI) frequencies, which are often many times higher than that of the redshifted 21-cm line (see table 2 of \\citealt{cmp+03}), and so in Fig. \\ref{2-distance-z} (bottom) we have flagged each of the DLAs according to how many times larger than $1420/(z_{\\rm abs}+1)$ the frequency at which the source size is measured, i.e. $\\nu_{\\rm size}/\\nu_{\\rm abs}$. The recently published low frequency VLBA imaging of \\citet{klm+09} accounts for many of the ``green'' values, where $\\nu_{\\rm size}/ \\nu_{\\rm abs}\\leq3$. Since a number of these sources appear resolved, we have recalculated the extents, rather than using the deconvolved sizes quoted in \\citet{klm+09}, which are often smaller than the beam and may be unphysical (e.g. $\\sim0$ pc for 0405--331).}. However, all of our non-detections (which are predominately at $z_{\\rm abs}\\gapp2.4$), also occult similar radio source sizes. Of these, three of the sizes are not as reliable (orange symbols where $3 < \\nu_{\\rm size}/ \\nu_{\\rm abs}<10$, cf. green where $\\nu_{\\rm size}/ \\nu_{\\rm abs}\\leq3$), although as seen from Fig. \\ref{Toverf}, only three published non-detections are sufficiently sensitive to detect $T_{\\rm spin}/{f}\\gapp2000$~K at $z_{\\rm abs}\\gapp2$ (Sect. \\ref{GMRT08}).  While, due to geometry effects, a high angular diameter distance will disadvantage the absorption cross-section in comparison to the same absorber placed at a low angular diameter distance, with the 21-cm searches spanning a look-back time range of 12 Gyr there are also evolutionary effects to be considered. For instance, an evolution in the morphologies of the absorbing galaxies is expected to result in different metallicities and these heavy elements can provide cooling pathways for the diffuse gas (\\citealt{dm72}). Furthermore, it is shown that the relative mix of the cold neutral medium (CNM, where $T\\sim150$ K and $n\\sim10$ \\ccm) and the warm neutral medium (WNM, $T\\sim8000$ K and $n\\sim0.2$ \\ccm) can be affected by the metallicity \\citep{whm+95}.  At low redshifts, where the DLA hosts can be imaged directly, the absorbing galaxies appear to be a mix of spirals, dwarf and LSBs (e.g. \\citealt{lbbd97,rnt+03,cl03}) and, in general, large galaxies appear to be more populous at low redshift with smaller morphologies becoming more common at high redshifts \\citep{bmce00,lf03}. Therefore, in the more compact absorbing galaxies, where the metallicities are lower than in the larger spirals (see Sect. \\ref{lsmc}), we may expect the gas to be more dominated by the WNM and, as shown by \\citet{ctp+07}, $T_{\\rm spin}/{f}$ is anti-correlated with the metallicity (see Sect. \\ref{lsmc}). However, although the metallicity is itself anti-correlated with the look-back time (\\citealt{pgw+03,rphw05}), no [M/H]--$z_{\\rm abs}$ correlation is seen for the 21-cm searched sample \\citep{ctp+07}, and, as seen in Fig.~\\ref{2-distance-z}, very few morphologies are known for these at $z_{\\rm abs}\\gapp1$. Furthermore, the $T_{\\rm spin}/{f}$\\,--\\,[M/H] anti-correlation is consistent with {\\em either} the spin temperature being lower or the covering factor being larger at high metallicities, with the larger galaxies likely to have high values of $f$ as well as [M/H]. Although the $T_{\\rm spin}/{f}$ degeneracy is unbreakable, it is possible that both factors contribute and that the values of [M/H], $T_{\\rm spin}$ and ${f}$ are in fact interwoven \\citep{ctp+07}.   In addition to decreasing metallicities, the increase in the background ultra-violet flux will increase the ionisation fraction of the gas \\citep{bdo88}, in addition to raising its spin temperature \\citep{fie59,be69}, an effect also seen in associated systems, where 21-cm absorption has never yet been detected in the host galaxy of a quasar with $L_{\\rm UV}\\gapp10^{23}$ W Hz$^{-1}$ \\citep{cww+08}. Therefore, as well as always having large angular diameter distance ratios and lower metallicities, high redshift absorption systems will generally be subject to higher ultra-violet fluxes, although for intervening absorbers the effect will not be as severe as for associated absorbers. Nevertheless, this could be the cause of many of the $T_{\\rm spin}/{f}$ lower limits at $z_{\\rm abs}\\gapp2.5$ (Fig. \\ref{Toverf}), although, again, such heated gas could have a significantly lower 21-cm absorption cross-section as well as a higher temperature.  \\subsubsection{Absorber extents} \\label{aex}  In order to consider the radio source size in conjunction with the ratio of the angular diameter distances, we calculate the extent of absorber required to fully cover the radio source (Fig.~\\ref{dla_schem}) and show this against the normalised line \\begin{figure} \\vspace{3.0cm} \\special{psfile=dla_schem.eps hoffset=-20 voffset=0 hscale=45 vscale=45 angle=0} \\caption{In the small angle approximation $\\theta = d_{\\rm abs}/DA_{\\rm DLA} = d_{\\rm QSO}/DA_{\\rm QSO}$, giving $d_{\\rm abs (min)}=\\theta_{\\rm QSO}.\\,DA_{\\rm DLA}$ when the absorber just covers the background radio source size.} \\label{dla_schem} \\end{figure} strength, $1.823\\times10^{18}\\,\\int\\tau\\,dv/N_{\\rm HI}= f/T_{\\rm spin}$ (Fig. \\ref{size-z}). \\begin{figure} \\vspace{6.5cm} \\special{psfile=size-foverT-klm+09-actual.ps  hoffset=-20 voffset=230 hscale=75 vscale=75 angle=270} \\caption{The minimum extent of the absorber required to cover the radio source versus the normalised line strength. The symbols are as per Fig. \\ref{2-distance-z} (bottom). The line shows the least-squares fit to the detections. The upper limits in the ordinate are due to upper limits in the background radio source sizes.} \\label{size-z} \\end{figure} From this we see a correlation between the extent of the absorber and the 21-cm absorption strength for the detections. For the 18 detections this is significant at $2.7\\sigma$, falling to $2.2\\sigma$ for the 12 detections with $\\nu_{\\rm size}/ \\nu_{\\rm abs}\\leq3$ (``green''), with the non-detections exhibiting no correlation whatsoever (bringing the significance for the whole 39 strong sample down to $1.4\\sigma$)\\footnote{Using the deconvolved values of \\citet{klm+09}, rather than those measured (see footnote \\ref{foot7}), gives $2.5\\,,1.9$ \\& $1.3\\sigma$, respectively.}.  The fit in Fig. \\ref{size-z} suggests that the non-detections which have been searched sufficiently deeply, (i.e. neglecting the two with $f/T_{\\rm spin}\\gapp0.01$), may require large absorbers ($\\gapp10$~kpc scales) in order to achieve effective covering factors. As is seen from the figure, several detections also occupy this regime, although these tend to be spiral galaxies.  Therefore, like \\citet{kc02}, and as discussed in \\citet{cmp+03}, the majority of non-detections (which are at high redshift and have unknown morphologies) may be non-spirals.  Although, unlike \\citet{kc02}, we again suggest that may be their small sizes, in addition to/rather than high mean spin temperatures, which may be responsible for their weak line strengths.  The covering factor can be defined\\footnote{It can also be estimated as the ratio of the compact unresolved component's flux to the total radio flux \\citep{bw83,klm+09}. However, even if high resolution radio images at the redshifted 21-cm frequencies are available, this method gives no information on the extent of the absorber (or how well it covers the emission). Furthermore, the high angular resolution images are of the continuum only and so do not give any information about the depth of the line when the extended continuum emission is resolved out.} as $f\\equiv [d_{\\rm abs}/(\\theta_{\\rm QSO}.\\,DA_{\\rm DLA})]^2$ \\citep{cw06}, meaning that, in general\\footnote{That is, where $d_{\\rm abs}$ is the actual extent of the absorber, rather than the extent minimum required to cover $\\theta_{\\rm QSO}$.}, the ordinate in Fig. \\ref{size-z} is equivalent to $d_{\\rm abs}/\\sqrt{f}$.  This therefore suggests, at least for the detections, that the line strength increases with the size of the 21-cm absorption region, a trend which is apparent, even when the absorber size is weighed down by the $\\sqrt{f}$ factor. and, in general, from the normalised line strength, \\begin{equation} 1.823\\times10^{18}\\,\\frac{\\int\\tau\\,dv}{N_{\\rm HI}}=  \\frac{f}{T_{\\rm spin}} \\propto\\theta_{\\rm QSO}.\\,DA_{\\rm DLA} = \\frac{d_{\\rm abs}}{\\sqrt{f}}. \\end{equation} That is, $d_{\\rm abs}\\propto f^{3/2}/T_{\\rm spin}$, which makes sense in terms of both covering factor and spin temperature -- a larger absorber implies a larger covering factor, as well as the possibility of a lower spin temperature. Presuming that this is not dominated by the numerator, this {\\em could} suggest that the larger absorbing galaxies are subject to lower spin temperatures (\\citealt{kc02} and references therein).  However, one interpretation of the distribution of the non-detections along the ordinate in Fig. \\ref{size-z} is that for a given absorption cross-section (i.e. assuming the same $d_{\\rm abs}$ for all the DLAs), the covering factor is generally lower than for a number of the detections.  Note that, in addition to geometry effects, it is possible that these lower covering factors could also be caused by lower beam filling factors, introduced by structure (``clumpiness'') in the absorbing gas, which may be behind the 21-cm variability seen at $z_{\\rm abs}=0.313$ towards 1127--145 \\citep{kc01} and at $z_{\\rm abs}=3.386$ towards 0201+113 \\citep{dob96,bbw97,kcl06}. \\citet{gsp+09a} also note from their large survey of Mg{\\sc \\,ii} absorption systems that the decreasing number density of 21-cm absorbers with redshift runs counter to that of the Lyman-\\AL\\ and ${\\rm W}_{\\rm r}^{\\lambda2796}>1$ \\AA\\ Mg{\\sc \\,ii} absorbers \\citep{rtn05}. This is attributed to a significant evolution in either the CNM filling factor or the covering factor, the former of which could be the result of a relatively extensive WNM, thus indicating a higher mean harmonic spin temperature \\citep{klm+09}, although the larger fractions of warm gas would not exclude the presence of cold clumps. Either or both of these scenarios would be tied up in the $T_{\\rm spin}/f$ degeneracy.  Nevertheless, 21-cm detection rates are significantly lower at high redshift and the filling factor of the CNM, through either its effect on $T_{\\rm spin}$ or $f$, could be a major contributor to this.  \\subsection{21-cm line strength--metallicity correlation} \\label{lsmc}  Addressing the 21-cm line strength--metallicity correlation ($T_{\\rm spin}/f$\\,--\\,[M/H] anti-correlation) of \\citet{ctp+07}, the focus of the 2008 GBT observations (Sect. \\ref{GBT08}), it appears that we should have detected absorption in at least one of the two DLAs for which reasonable sensitivities were reached (Fig. \\ref{foverT-m}). \\begin{figure} \\vspace{6.5cm} \\special{psfile=foverT-m.ps  hoffset=-20 voffset=230 hscale=75 vscale=75 angle=270} \\caption{The normalised line strength ($\\equiv f/T_{\\rm spin}$) versus the metallicity for the DLAs searched in 21-cm absorption, where the symbols are as per Fig.~\\ref{Toverf}. The significance of the correlation is given for both the whole sample and the detections only and the least-squares fit for the detections shown.} \\label{foverT-m} \\end{figure} Regarding the best improved limit\\footnote{By a factor of $\\approx7$, cf. \\citealt{cld+96}. See footnote \\ref{cld}.}, 21-cm absorption may not have been detected towards 0528--250, on the grounds that, with an absorption redshift very close to that of the emission redshift (both at $z=2.8110$, \\citealt{mff04} and references therein), the absorber is possibly located in the host of the background quasar. We may therefore expect the neutral gas to be subject to excitation effects seen in other high redshift 21-cm absorption searches \\citep{cww+08}. However, the presence of H$_2$ in the absorber suggests that the gas is relatively cool with $T_{\\rm kin}\\approx 110 - 150$ K (\\citealt{spl+05}) and so this is unlikely.  In any case, referring to Fig.~\\ref{foverT-m}, the limit is not overly low in comparison to 0438--436 (the detection at [M/H]$\\,=-0.68$, $f/T_{\\rm spin}=6.3\\times10^{-4}$ [K$^{-1}$]) and there is also the possibility that we have missed the line, with our limit only being good over the redshift range $z\\approx 2.806-2.812$ (Sect. \\ref{GBT08}).  Assuming this limit is reliable and including 0311+430, strengthens the correlation (cf. \\citealt{ctp+07}).  \\citet{ykep07} comment that the $z_{\\rm abs}=2.289$ absorber towards 0311+430 has an unusually high metallicity, although according to Fig. \\ref{foverT-m}, this has some leeway before the metallicity becomes atypically high. From its grouping in the line strength--metallicity plot, we suggest that the absorption is due to a spiral galaxy, as do \\citet{ykep07} on the basis of its low ``spin temperature''. Support for this assertion is given by its grouping with the spirals in Fig. \\ref{size-z} (the ``unknown'' at $7\\times10^{-3}$ K$^{-1}$, 5 kpc), although this is on the grounds of a large absorption cross-section, rather than the spin temperature."
    },
    "0910/0910.5813_arXiv.txt": {
        "abstract": "{}{The detection of rotational transitions of the AlO radical at millimeter wavelengths from an astronomical source has recently been reported. In view of this, rotational transitions in the ground \\textit{X}\\textsuperscript{2}$\\Sigma$\\textsuperscript{+} state of AlO have been reinvestigated. } {Comparisons between Fourier transform and microwave data indicate a discrepancy regarding the derived value of $\\gamma$$_{D}$ in the v = 0 level of the ground state. This discrepancy is discussed in the light of comparisons between experimental data and synthesized rotational spectra in the v = 0, 1 and 2 levels of \\textit{X}\\textsuperscript{2}$\\Sigma$\\textsuperscript{+}.} {A list of calculated rotational lines in v = 0, 1 and 2 of the ground state up to \\textit{N}' = 11 is presented which should aid astronomers in analysis and interpretation of observed AlO data and also facilitate future searches for this radical.}{} ",
        "introduction": "\\paragraph{ The ground \\textit{X}\\textsuperscript{2}$\\Sigma$\\textsuperscript{+} state of the AlO radical has been studied with microwave spectroscopy in the v = 0, 1 and 2 levels. T\\\"{o}rring et al. (1989) recorded the \\textit{N} = 1 $\\rightarrow$2 transition near 76 GHz. In their Table 1, frequencies for several $\\Delta$\\textit{F} = $\\Delta$\\textit{N} and $\\Delta$\\textit{F} $\\neq$ $\\Delta$\\textit{N} transitions are shown. Yamada et al. (1990) and Goto et al. (1994) recorded several $\\Delta$\\textit{F} = $\\Delta$\\textit{N} rotational transitions, resulting in a set of accurate molecular constants, including $\\gamma$ and $\\gamma$$_{D}$. } \\paragraph{  Launila et al. (1994) have performed a Fourier transform study of the \\textit{A}\\textsuperscript{2}$\\Pi$$_{i}$\\textbf{$\\rightarrow$}\\textit{X}\\textsuperscript{2}$\\Sigma$\\textsuperscript{+} transition of AlO in the 2 $\\mu$m region. In that work, some discrepancies were pointed out regarding their derived $\\gamma$$_{D}$ values, as compared to those found in the microwave work. While Yamada et al. (1990) had given a positive value for $\\gamma$$_{D}$ for v = 0, the sign was in fact found to be negative in the light of high-\\textit{N} data of Launila et al. (1994). One of the aims of the present paper is to reinvestigate this discrepancy more closely. The work by Goto et al. (1994), dealing with the v = 1,2 vibrational levels of the ground state of AlO, does not show the same discrepancy, however. } \\paragraph{In a theoretical work, Ito et al. (1994) have discussed and explained the observed vibrational anomalies in the spin-rotation constants of the ground state in the light of spin-orbit interaction with the \\textit{A}\\textsuperscript{2}$\\Pi$$_{i}$ and \\textit{C}\\textsuperscript{2}$\\Pi$ states.} \\paragraph{Recently, Tenenbaum and Ziurys (2009) reported three rotational transitions of AlO from the supergiant star VY Canis Majoris. In order to facilitate future millimeter-wave search for rotational transitions of AlO, tables containing expected frequencies in the v = 0, 1 and 2 levels of the ground state are useful. In the present work, such tables are presented.} ",
        "conclusions": "\\paragraph{The determination of the correct constants for a molecule/radical has its own intrinsic value. Additionally, the present study is also relevant in an astronomical context. Tenenbaum and Ziurys (2009) have recently made the first radio/mm detection of the  AlO (\\textit{X}\\textsuperscript{2}$\\Sigma$\\textsuperscript{+}) radical toward the envelope of the oxygen rich supergiant star VY Canis Majoris (VY CMa). They observed the \\textit{N} = 7$\\rightarrow$6 and 6$\\rightarrow$5 rotational transitions of AlO at 268 and 230 GHz and the N = 4\\textbf{$\\rightarrow$}3 line at 153 GHz. While their search for the \\textit{N} = 7$\\rightarrow$6 hyperfine transitions was based on direct laboratory measured frequencies by Yamada et al  (1990), the search for the \\textit{N} = 6$\\rightarrow$5 and \\textit{N} = 4$\\rightarrow$3 transitions was based on frequencies calculated from spectroscopic constants of those authors. As has been pointed out, an error is present in one of these constants ($\\gamma$$_{D}$ in the v = 0 level of the ground state), which leads to deviations between the calculated and true frequencies. While these deviations are small and may not affect the final result of a line search too seriously, it is still desirable to have accurate frequencies to  facilitate future millimeter-wave searches for rotational transitions of AlO. It is planned for several such searches to be taken place in the near future as a consequence of the recent detection in VY CMa. AlO is slowly emerging as a molecule which could attract a fair deal of interest among astronomers. In the VY CMa detection for example it is shown how its study could lead to a better understanding of the gas-phase refractory chemistry in oxygen-rich envelopes. AlO has also been proposed to be a potential molecule in the formation of alumina, which is one of the earliest and most vital dust condensates in oxygen rich circumstellar environments (Banerjee et al. 2007). The mineralogical dust condensation sequence and processes involved are issues of considerable interest to astronomers. The AlO radical also received considerable attention after the strong detection of several \\textit{A}$\\rightarrow$\\textit{X} bands in the near-infrared (1 - 2.5 microns) in the eruptive variables V4332 Sgr and V838 Mon (Banerjee et al. 2003; Evans et al 2003). Other IR detections of AlO include IRAS 08182-6000 and IRAS 18530+0817 (Walker et al. 1997). In the optical, the \\textit{B}$\\rightarrow$\\textit{X} bands have been prominently detected in U Equulei (Barnbaum et al. 1996) and in several cool stars and Mira variables including Mira itself (Keenan et al. 1969;  Garrison 1997).} \\paragraph{The line lists presented here should also aid in analysis of mm line data for aspects related to kinematics. Since the rotational levels of AlO species are split by both fine and hyperfine interactions, each rotational transition consists of several closely spaced hyperfine components. Figure 2 exemplifies this. If line broadening is small and the spectral resolution of the observations is adequate to resolve these components, then different components will yield different Doppler velocities if the corresponding reference or rest frequencies are in error.  This could lead to ambiguity in interpreting the data. Even if the hyperfine components are not distinctly resolved but rather blended to give a composite line profile (as in the observed profiles in VY CMa), modelling of such composite profiles using wrong rest frequencies could lead to errors in estimating  the composite line centre, half width of the composite profile and half-widths of the individual hyperfine components. Proper estimates of such kinematic parameters are important as they help determine the size and site of origin of a molecular species. The detailed discussion by Tenenbaum and Ziurys (2009) in the case of VY CMa, which has three distinctly different kinematic flows in the system, illustrates this.}"
    },
    "0910/0910.3233_arXiv.txt": {
        "abstract": "The detection of upward propagating internal gravity waves {\\glite at the base of} the Sun's chromosphere has recently been reported by Straus et al., who postulated that these may efficiently couple to Alfv\\'en waves in magnetic regions. This may be important in transporting energy to higher levels. Here we explore the propagation, reflection and mode conversion of linear gravity waves in a VAL C atmosphere, and find that even weak magnetic fields usually reflect gravity waves back downward as slow magnetoacoustic waves well before they reach the Alfv\\'en/acoustic equipartition height at which mode conversion might occur. However, for certain highly inclined magnetic field orientations in which the gravity waves manage to penetrate near or through the equipartition level, there can be substantial conversion to either or both upgoing Alfv\\'en and acoustic waves. Wave energy fluxes comparable to the chromospheric radiative losses are expected. ",
        "introduction": "\\label{intro} Using the Interferometric BIdimensional Spectrometer (IBIS) and the Echelle Spectrograph on the Dunn Solar Telescope (DST) of the Sacramento Peak National Solar Observatory, and the Michelson Doppler Imager (MDI) on the Solar and Heliospheric Observatory (SOHO), \\cite{straus08} have identified upward propagating\\footnote{The group velocity and hence energy flux is upward. As expected of gravity waves, the phase velocity is downward.}  gravity waves with frequencies between 0.7 mHz and 2.1 mHz in weak magnetic field regions of the solar atmosphere, and found that their energy flux is an order of magnitude larger than co-spatial acoustic waves, and comparable to the expected quiet-sun chromospheric losses of around 4.3 $\\rm kW\\,m^{-2}$. However, the gravity waves were found to be significantly suppressed in stronger field regions. Nevertheless, in light of recent identification of ubiquitous Alfv\\'en waves in the corona \\citep{pontieu07,tomczyk07}, they postulate that when these gravity waves enter magnetic regions they may efficiently couple to Alfv\\'en waves, perhaps contributing to the observed coronal wave flux. Furthermore, \\cite{jess09} have directly identified torsional Alfv\\'en waves in H$\\alpha$ bright-point groups at frequencies as low as 1.4 mHz, which places them at least partially in the gravity wave regime if there is any coupling.  In this paper, we explore the propagation, reflection, and mode conversion of atmospheric gravity waves of around 1 mHz in frequency using both dispersion relations and numerical solution of the governing differential equations in simple atmospheric models with uniform inclined magnetic field. The imposed fields of 10 to 100 Gauss are weak in the photospheric context, but become dominant at greater heights as the plasma $\\beta$ (the ratio of plasma to magnetic pressure) falls below unity due to density stratification. We therefore have a situation where low frequency waves are essentially gravity waves at low altitudes, but become magnetically dominated at higher levels in the chromosphere. The central question is: \\emph{What happens to upward propagating gravity waves as they enter regions where magnetic forces become significant? Does the magnetic field help or hinder propagation through the chromosphere?} The answer is: both, depending on magnetic field orientation. ",
        "conclusions": "\\label{conclusions} The main conclusions we draw from our analyses are: \\renewcommand{\\labelenumi}{\\arabic{enumi}.} \\begin{enumerate} \\item Even very weak magnetic fields very effectively reflect gravity waves back downward as slow magneto-acoustic waves. This typically happens well below the $a=c$ equipartition level. \\item However, at very large magnetic field inclinations, typically around $80^\\circ$ or more depending on frequency (see Fig.~\\ref{fig:Ffreqs}), substantial mode conversion from gravity waves to either field-guided acoustic waves (for small $\\phi$) or Alfv\\'en waves ($20^\\circ\\la \\phi \\la 70^\\circ$) occurs, and these waves continue to propagate upward along the field lines. The amount of energy they carry is potentially significant for the upper chromosphere. \\item Wave energy fluxes reaching the top of our model are very sensitive to magnetic field direction, but quite insensitive to magnetic field strength, at least in the range 10 G -- 100 G. \\item The dispersion diagrams give a simple, easy and quite accurate picture of the behaviour of gravity waves in a magneto-atmosphere, though with the caveat that tunnelling can sometimes occur between branches. \\end{enumerate}  In simple terms, we conclude that atmospheric gravity waves are very effectively suppressed by even very weak magnetic field, \\emph{unless} that field is highly inclined and the attack angle fine, in which case it opens a window to the upper atmosphere that allows the gravity waves to propagate through in a different guise. This is closely related to the ``magnetic portals'' of \\cite{jeff06} which rely on the ramp effect to allow acoustic waves below the acoustic cutoff frequency to still propagate upward in low-$\\beta$ inclined magnetic field, but in the case of low-frequency gravity waves which are already propagating, it gives them the opportunity to convert to propagating acoustic waves around the $a=c$ level. Because of their low frequency though, very substantial inclination is required to open these windows, but no more than is characteristic of chromospheric canopy.  {\\hilite There are consequences for recent and future observations of solar atmospheric oscillations. The surprising extent to which even very weak vertical or moderately inclined magnetic field inhibits gravity waves by causing them to quickly reflect as slow magneto-acoustic waves perhaps explains ``significantly suppressed atmospheric gravity waves at locations of magnetic flux'' found by \\cite{straus08}. The ubiquity of near-horizontal field in the low solar atmosphere discovered recently with \\emph{Hinode} \\citep{lites08} however raises the possibility that low frequency gravity waves may efficiently couple to Alfv\\'en waves that continue to propagate vertically into the corona, contributing to the vast sea of waving field lines now known to exist there \\citep{pontieu07,tomczyk07}. Although \\citeauthor{tomczyk07} detect a peak in velocity power of these Alfv\\'en oscillations at around 3.5 mHz, (presumably driven by the Sun's internal normal modes), the power spectrum continues to rise with decreasing frequency till at least 1 mHz, well inside the gravity wave regime at photospheric level. It is tempting to postulate that gravity waves may be the vector of this wave energy at low levels and that it may convert to Alfv\\'en waves around the acoustic/Alfv\\'enic equipartition level in highly inclined field regions. Further observational work, ideally at multiple heights, is warranted to more fully explore these possibilities. }  It should be emphasised though that our analysis is entirely linear. Acoustic waves are likely to shock before reaching the upper chromosphere. However, Alfv\\'en waves do not suffer this fate. Our models are also adiabatic, which is not a good representation of the chromosphere, {\\hilite though the detections of \\cite{straus08} suggest that {\\glite atmospheric} radiative losses do not completely suppress gravity waves, at least at the {\\glite photospheric} altitudes sampled by IBIS} {\\glite and MDI}. The adiabatic assumption will be relaxed in future work."
    },
    "0910/0910.3834_arXiv.txt": {
        "abstract": "The bulk motion of galaxies induced by the growth of cosmic structure offers a rare opportunity to test the validity of general relativity across cosmological scales.  However, modified gravity can be degenerate in its effect with the unknown values of cosmological parameters. More seriously, even the `observed' value of the RSD (redshift-space distortions) used to measure the fluctuation growth rate depends on the assumed cosmological parameters (the Alcock-Paczynski effect).  We give a full analysis of these issues, showing how to combine RSD with BAO (baryon acoustic oscillations) and CMB (Cosmic Microwave Background) data, in order to obtain joint constraints on deviations from general relativity and on the equation of state of dark energy whilst allowing for factors such as non-zero curvature.  In particular we note that the evolution of $\\Omega_m(z)$, along with the Alcock-Paczynski effect, produces a degeneracy between the equation of state $w$ and the modified growth parameter $\\gamma$. Typically, the total marginalized error on either of these parameters will be larger by a factor $\\simeq 2$ compared to the conditional error where one or other is held fixed. We argue that future missions should be judged by their Figure of Merit as defined in the $w_p - \\gamma$ plane, and note that the inclusion of spatial curvature can degrade this value by an order of magnitude. ",
        "introduction": "Models of dark energy leave a characteristic signature embedded in both the cosmic expansion and structure formation histories. Recent observational progress has been made with the former, due to its relative ease of measurement, leading to a measurement of the dark energy equation of state parameter $w\\equiv P/\\rho c^2$ with better than 10\\% precision \\cite{2008arXiv0803.0547K,2009arXiv0907.1660P}. This work is geometrical, and so probes dark energy only through its influence on the evolving expansion rate of the Universe. It is thus possible that dark energy may be an illusion, indicating the need to revise general relativity and thus also the Friedmann equation.  In either case, the phenomenological dark energy term may well differ from a cosmological constant ($w=-1$), and may change its equation of state with redshift. These possible degrees of freedom need to be allowed for before we can claim any evidence for a deviation from general relativity. This paper thus considers how we can make simultaneous measurements of the properties of dark energy and of modified gravity.   A number of probes are capable of measuring $w$ via its influence on the redshift-distance relation. This measurement alone is effectively completely degenerate with a modification of gravity on the scale of the Hubble radius. But for many models, the Mpc scales of galaxy clustering may be affected in a different way; the growth rate of density fluctuations has thus emerged as a key means of breaking this degeneracy between gravity and dark energy \\cite{2005PhRvD..72d3529L}\\cite{2008Natur.451..541G}. It is rather more difficult to study the growth rate, due to uncertainty in the behaviour of galaxy bias, but there are currently two promising avenues available for future exploration. Weak gravitational lensing provides a direct measurement of the dark matter distribution, and its evolution with redshift. It can also probe broader aspects of modified gravity, particularly the balance between perturbations to the time and space parts of the metric \\cite{beanism,2008PhRvD..78f3503J}. The focus of the present work will be the alternative technique, known as redshift-space distortions (RSD), which exploit the relationship between the large-scale coherent velocities of galaxies and the growth rate of perturbations.  In real space, we expect the clustering of galaxies to be statistically isotropic. However, in redshift space the line-of-sight component of a galaxy's peculiar velocity breaks this symmetry. Inside a virialized cluster of galaxies, the orbital velocity dispersion scatters galaxy redshifts, creating the `Fingers of God', and thereby erasing spatial information on small scales. Across larger scales, galaxies coherently fall out of voids and into overdense regions, considerably amplifying the power in redshift space. These two effects are often treated independently, although a more complex model is required to attain a higher degree of precision \\cite{scocciz}. For the present purpose, the large-scale effect is the aspect of interest, since continuity relates coherent peculiar velocities directly to the growth rate of density fluctuations.  Observations to date have led to estimates of the growth rate at various redshifts, although not yet at a useful level of precision \\cite{2002MNRAS.332..311H,2005MNRAS.361..879D,2007MNRAS.381..573R,2008Natur.451..541G}. Future surveys are likely to cover orders of magnitude larger volumes, thereby delivering the precision needed to discriminate interesting models of modified gravity. But we shall see that, when approaching this target, it may no longer be appropriate to make the simplifying assumptions adopted to date.  In \\S\\ref{sec:sig} we review the process of determining the growth rate from redshift distortions, before constructing a Fisher matrix. Our results are presented in \\S\\ref{sec:results}, while in \\S\\ref{sec:merit} we consider the implications for the proposed dark energy Figure of Merit. ",
        "conclusions": "By relaxing the common assumption of a fixed background cosmology, we have highlighted some of the difficulties encountered when attempting to study gravity via the bulk motion of galaxies. Rather than a pure probe of structure, redshift distortions also comprise a geometric component. This enters at the stage of converting the true observables, angles and redshifts, into distances and Fourier modes. Furthermore, when determining the growth index $\\gamma$ it is essential that its corresponding radix $\\Omega_m(z)$ is well determined. With these two factors in mind, it appears unlikely that the galaxy power spectrum alone could provide conclusive evidence against General Relativity.  To converge on the true underlying cosmology, iterating over a value for $\\Omega_m$ has proved adequate for current data. However with the greater degrees of freedom required to test relativity ($w_0, w_a, \\Omega_k$), the available volume of parameter space appears too great. Fortunately future data will inevitably be accompanied by improved measurements of the baryon acoustic oscillations. Ironically the squashing effect that empowers the BAO is the very same Alcock-Paczynski effect that confounds the redshift distortions.  One concern in the formalism may be the assumption of scale-independence for both the growth and bias. More physically motivated forms of modified gravity, such as $f(R)$ models, \\cite{2007PhRvD..76j4043H}, lead to rather different scale-dependent growth factors. However, as highlighted in \\cite{2007PhRvD..76j4043H}, such models also generate very prominent deviations on intermediate scales, which would become more immediately apparent.  Nevertheless, neglect of these issues is more likely to lead to a bias in the results of analyses that assume scale-independent effects, rather than changing their statistical precision. In this work, we have concentrated on the latter aspect, and our main conclusion is that the parameters $\\gamma$ and $w_p$ will generally be strongly anti-correlated. We therefore suggest that a natural Figure of Merit for future experiments in fundamental cosmology should be the reciprocal of the area of the error contour in the $\\gamma - w_p$ plane.    \\noindent{\\bf Acknowledgements} \\\\ We thank Luigi Guzzo for many helpful comments on an earlier draft of this paper, and also Thomas Kitching and Will Percival for several productive discussions. FS was supported by an STFC Rolling Grant."
    },
    "0910/0910.0837_arXiv.txt": {
        "abstract": "We examine the star formation rates (SFRs) of galaxies in a redshift slice encompassing the $z = 0.834$ cluster \\rxjfull.  We used a low-dispersion prism in the Inamori Magellan Areal Camera and Spectrograph (IMACS) to identify galaxies with $z_{\\rm AB} < 23.3$~mag in diverse environments around the cluster out to projected distances of $\\sim 8$~Mpc from the cluster center.  We utilize a mass-limited sample (\\mcutrange) of \\nmcut\\ galaxies that were imaged by Spitzer MIPS at 24~\\micron\\ to derive SFRs and study the dependence of specific SFR (SSFR) on stellar mass and environment.  We find that the SFR and SSFR show a strong decrease with increasing local density, similar to the relation at $z \\sim 0$.  Our result contrasts with other work at $z \\sim 1$ that find the SFR-density trend to reverse for luminosity-limited samples. These other results appear to be driven by star-formation in lower mass systems ($M \\sim 10^{10}$~\\msun).  Our results imply that the processes that shut down star-formation are present in groups and other dense regions in the field.  Our data also suggest that the lower SFRs of galaxies in higher density environments may reflect a change in the ratio of star-forming to non-star-forming galaxies, rather than a change in SFRs.  As a consequence, the SFRs of star-forming galaxies, in environments ranging from small groups to clusters, appear to be similar and largely unaffected by the local processes that truncate star-formation at $z \\sim 0.8$. ",
        "introduction": "Pioneering studies of the impact of environment on galaxy properties have found higher density environments in the local universe to be dominated by elliptical and S0 galaxies \\citep{dressler1980}.  Because of the strong correlation of Hubble type with star formation rate \\citep[SFR,][]{kennicutt1998}, a correlation between SFR and local galaxy density is expected and has been seen in more recent studies of the local universe \\citep{gomez2003,balogh2004}.  \\citet{kauffmann2004} examined the SDSS sample of \\citet{brinchmann2004} and found a strong dependence of SFR on local galaxy density and stellar mass at $z \\sim 0$.  At a fixed stellar mass, they found the specific SFRs (SSFRs), and therefore SFRs, of galaxies to be lower in higher density environments.  At $z \\sim 1$, the morphology-density relation (MDR) follows a similar trend to that found in the local universe for mass-limited samples \\citep{holden2007,vanderwel2007b}.  In contrast, recent analyses of the field at $z \\sim 1$ suggest that the SFR-density trend was the {\\em reverse} at earlier times, where galaxies at higher densities display {\\em higher} SFRs than galaxies at lower densities \\citep{elbaz2007,cooper2008}.  In this Letter, we explore the SFRs of galaxies in a redshift slice that includes the $z = 0.834$ galaxy cluster \\rxjfull\\ (hereafter \\rxj).  Our wide-field spectroscopic survey extends to a projected distance of $\\sim 8$~Mpc from the cluster center, resulting in a large number of galaxies that span a much broader range of environments at this epoch than is found in other work.  Our aim is to determine the form of the SSFR/SFR-density relation at $z \\sim 0.8$ and compare to the results discussed above at these redshifts and at $z \\sim 0$.  We assume a cosmology with $H_0 = 70$~\\kmps~\\pmpc, $\\Omega_M = 0.3$, and $\\Omega_{\\Lambda} = 0.7$.  Stellar masses and SFRs are based on a Chabrier IMF \\citep{chabrier2003}. ",
        "conclusions": "Recent work at $z \\sim 1$ indicates that for galaxies in the mass range studied here, several relations follow trends similar to those found at $z \\sim 0$.  For example, \\citet{holden2007} and \\citet{vanderwel2007b} found a strong MDR for mass-limited samples in both the field and clusters at $z \\sim 1$.  Thus, it should not be surprising that we see a similar trend for the SSFR/SFR-density relation at $z \\sim 0.8$ as we see at $z \\sim 0$ (although with a different normalization).  Interestingly, some recent studies of galaxies at low-redshift find evidence for enhanced levels of dust-obscured SF at densities slightly above typical field densities, although they also generally find an overall trend of decreasing SF activity at higher densities \\citep{gallazzi2009,wolf2009,haines2009b}.   Our result of a declining SFR in higher density environments appears universal at $z \\sim 0.8$ and not confined to a cluster environment, much like the MDR and SFR-density relations in the local universe.  As seen in Figure~\\ref{stack_mips_density}, cluster galaxies dominate the highest density bin, but much of the decrease in SSFR occurs in lower density field bins.  In addition, after removing galaxies within $\\sim 2.5$ times the virial radius of the cluster, we continue to find the SSFR to decrease with density, including in the remaining high density regions represented by several groups at projected distances of $\\sim 3-5$~Mpc from the cluster.  This decrease in SSFR occurs over a similar range of densities in which \\citet{patel2009} found an increase in the red galaxy fraction, possibly linking the end of SF and buildup of red-sequence galaxies in environments that reach into the dense regions of the field.   In contrast to our work, \\citet{elbaz2007} and \\citet{cooper2008} found very different results for the SFR-density relation.  At $0.8 < z < 1.2$, \\citet{elbaz2007} found a factor of $\\sim 6$ spread in SFRs.  However, they found galaxies at higher densities to have {\\em higher} mean SFRs up to a critical density, above which the SFR declined.  Much of the reversed SFR-density trend in \\citet{elbaz2007} is driven by a peak in the SFR in a narrow projected density range of $\\sim 0.1$~dex ($3 < \\Sigma~(\\rm Mpc^{-2}) < 4$).  Likewise, \\citet{cooper2008} also found the SFRs of galaxies to increase in higher density environments at $0.75 < z < 1.05$, although their observed spread in SFRs was less than a factor of $\\sim 1.5$.  While neither of these two surveys contain a cluster, both sample group environments similar to the groups around \\rxj\\ that have velocity dispersions \\citep[$\\sim 400$~\\kmps,][]{tanaka2006} that are typical of groups found in the field \\citep{gerke2007}.  However, the reversal in the SFR-density relation found by \\citet{elbaz2007} and \\citet{cooper2008} does not extend into the highest densities found in their group environments.  Sample selection plays an important role in interpreting the different results.  While we use a mass-limited sample, \\citet{elbaz2007} and \\citet{cooper2008} use luminosity-limited samples that are biased to include low-mass blue star-forming galaxies and exclude the corresponding non-star-forming red galaxies.  When restricted to a mass-limited sample that was similar to the one in this work, \\citet{elbaz2007} found the SSFR continued to increase from low-density to high-density but with marginal significance (see Fig.~20 in that work).  \\citet{cooper2008} found the mean SFR to increase with density by $\\sim 35\\%$ while for the same sample the mean SSFR {\\em decreases} by $\\sim 35\\%$ over the same density range, which implies an increase in the mean mass with density by a factor of $\\sim 2$.  The inferred range of mean mass ($9.9 \\la \\log M/M_{\\odot} \\la 10.2$) from \\citet{cooper2008} is below our mass threshold.  However, the small rise in SFRs with density seen by \\citet{cooper2008} can be explained by the presence of relatively more high mass galaxies, which have higher SFRs \\citep{elbaz2007,noeske2007b}, at higher densities in their sample.   Interestingly, we note that when selecting a luminosity-limited sample that was similar to either of these works, we found the SFR to continue to decrease at higher densities, but at a much shallower pace for lower mass galaxies.  This did not depend strongly on the sub-sample of galaxies that were used to compute local densities (i.e. luminosity vs. mass-limited).  We note that contaminants from low-density environments, which we found to have higher levels of SF, likely contribute to the shallower trend.   Recent observations of a ``star-forming sequence'' in the field at $z \\sim 0$ \\citep{salim2007} and up to redshifts of $z \\sim 1$ \\citep{noeske2007} suggest that galaxies populate a narrow distribution of SSFRs at a fixed stellar mass.  Below this sequence is a more extended distribution of galaxies with very low SSFRs that are effectively non-star-forming.  Here, we investigate whether the SSFR-density trend represents (1) a changing mix of galaxies that lie on or off of the star-forming sequence or (2) a change in the SFRs of star-forming galaxies in different environments.  Without additional constraints, our MIPS stacking analysis alone cannot distinguish between these two scenarios.  Using our previous results on the correlation between the red galaxy fraction and density \\citep{patel2009}, we can speculate on which of these scenarios is more likely.  In \\citet{patel2009} we found the fraction of blue galaxies increased by a factor of $\\sim 10$ from $\\sim 5\\%$ in the highest density regions to $\\sim 50\\%$ at the lowest densities.  If the typical SFR of galaxies on the star-forming sequence remains constant and the ratio of blue-to-red galaxies reflects the ratio of galaxies that are forming stars or not then one expects the average SFR (and SSFR within a single mass bin) in the high density regions to be diminished by a factor of $\\sim 10$ compared to the field, roughly what we see in the MIPS stacking analysis presented here.  The notion that galaxies are either star-forming or not is given additional credence by \\citet{bell2005} and \\citet{dressler2009b} who found that the dominant mode of SF at these epochs was in moderate starbursts when accounting for the duty cycle of such events.  In summary, we find that at $z \\sim 0.8$, the SSFRs and SFRs of galaxies with \\mcutrange\\ decrease in higher density environments, and this is true over the entire range of environments studied.  This result follows the trend observed at $z \\sim 0$ in which \\citet{kauffmann2004} also found galaxies at higher densities to have lower SSFRs, but with SSFRs that were $\\sim 10$ times higher at $z \\sim 0.8$.  Our result differs from that of \\citet{elbaz2007} and \\citet{cooper2008} who found a reversal in the SFR-density trend at $z \\sim 1$ over some portion of their density range, mostly driven by SF in lower mass galaxies present in their luminosity-limited samples.  In a forthcoming paper, we plan to utilize rest-frame UV data with SED fitting to determine total SFRs of individual galaxies at $z \\sim 0.8$.  The distribution of SFRs in different environments will allow us to probe the processes that were responsible for shutting down SF at a time when the universe was half its current age."
    },
    "0910/0910.0006_arXiv.txt": {
        "abstract": " ",
        "introduction": "We observed a  field centered on the asynchronous polar CD Ind (magnetic cataclysmic binary; see \\citeauthor{CDInd1}, \\citeyear{CDInd1} and \\citeauthor{CDInd2}, \\citeyear{CDInd2}) from December 21, 2008 until January 19, 2009 during 10 nights. Observations were done remotely at Tzec Maun Observatory (Pingelly, Western Australia) using the Takahashi TOA-150 apochromatic refractor (D=150 mm, F = 1095 mm) and SBIG STL-6303 CCD camera. All observations were conducted with the Bessel R filter. The field of view was $87' \\times 58'$. About 4 000 objects were detected on each frame. Our goal was to find new variable stars. ",
        "conclusions": "We examined light curves of 3712 objects in the $87' \\times 58'$ field centered on CD Ind and discovered three new variable stars. Two of them are most probably  W UMa-type eclipsing binaries (EW), and the other one is most probably an  RR Lyr-type (RRC) star. We show that \\emph{RMS-scatter vs. Magnitude} diagram is a useful tool to discover new variable stars.  Table \\ref{table:summary} summarizes parameters of discovered variable stars, such as USNO-B1.0 designation, equatorial coordinates, brightness in maximum and minimum, period (in fractions of days), zero-epoch ($E_0$, time of maximum for RRC star and time of minimum for EW stars) and GCVS type of variability."
    },
    "0910/0910.3763_arXiv.txt": {
        "abstract": "We derive kinematic distances to the 86 6.7 GHz methanol masers discovered in the Arecibo Methanol Maser Galactic Plane Survey. The systemic velocities of the sources were derived from \\coii~($J=2-1$), CS ($J=5-4$), and \\nhiii~observations made with the ARO Submillimeter Telescope, the APEX telescope, and the Effelsberg 100 m telescope, respectively. Kinematic distance ambiguities were resolved using \\hi~self-absorption with \\hi~data from the VLA Galactic Plane Survey. We observe roughly three times as many sources at the far distance compared to the near distance. The vertical distribution of the sources has a scale height of $\\sim 30$ pc, and is much lower than that of the Galactic thin disk. We use the distances derived in this work to determine the luminosity function of 6.7 GHz maser emission. The luminosity function has a peak at approximately $10^{-6}~L_\\sun$. Assuming that this luminosity function applies, the methanol maser population in the Large Magellanic Cloud and M33 is at least 4 and 14 times smaller, respectively, than in our Galaxy. ",
        "introduction": "Methanol masers at 6.7 GHz are unique compared to their OH and H$_2$O counterparts in that they appear to be exclusively associated with early phases of massive star formation. They are hence extremely useful tools to identify and study very young massive star forming regions. To date, over 800 sources have been detected through various targeted and blind surveys of the Galactic plane (e.g. \\citealt{gree09}; compilation of \\citealt{xu09b}).  In spite of the extensive surveys to date for 6.7 GHz methanol masers, their luminosity function remains largely unknown. Since 6.7 GHz methanol masers are closely associated with mainline OH masers in massive star forming regions, \\citet{casw95} suggested that the luminosity function of methanol masers would be similar to that of OH masers, but with a scaling factor to account for the fact that methanol masers are on average much more intense than OH masers. \\citet{van96} use a probabilistic approach to assign distances and estimated the luminosity function to have a power-law behavior with an index around --2. Both these studies used only the peak flux density rather than the integrated flux for determining the luminosity functions. The most recent study of the issue was carried out by \\citet{pest07} who used a compilation of all sources detected prior to 2005 \\citep{pest05}. Carrying out a statistical analysis of the maser population and modeling the spatial distribution of the masers, the luminosity function was modeled by those authors as a power-law with sharp cutoffs and having an index between --1.5 and --2.  There are two problems facing studies of the luminosity function. First, most studies use a catalog of methanol masers compiled from various surveys, many of them targeted (towards IRAS sources or OH masers), and with different sensitivities. However, unbiased searches have all shown that targeted searches, especially towards IRAS sources, underestimate the number of sources by a factor of 2 or greater \\citep{pand07b,szym02,elli96}. For the best estimate of the luminosity function, one should employ a blind survey, since it is possible to analyze the limitations such as completeness reliably. Among the several blind surveys to date, by far the most sensitive one is the Arecibo Methanol Maser Galactic Plane Survey (AMGPS; \\citealt{pand07a}). The AMGPS covered an area of 18.2 square degrees between Galactic longitudes of 35\\degr~and 54\\degr, and detected a total of 86 sources. This survey has a 95\\% probability of detection at a peak flux density of 0.27 Jy although sources as weak as 0.13 Jy have been detected. This makes the AMGPS catalog ideal for determining the luminosity function of 6.7 GHz methanol masers, especially at faint luminosities, although the relatively small size of the sample and area covered by the survey results in large statistical uncertainties.  A more formidable problem is the determination of distances to the sources. Measuring distances is an old and challenging problem in observational astrophysics. While trigonometric parallax is the most reliable method for determining distances, it cannot be applied readily to a large sample of sources. The usual technique for Galactic sources is to use a Galactic rotation curve to determine kinematic distances. While kinematic distances can have significant errors at times (e.g. \\citealt{xu06}), recent work using Very Large Baseline Interferometry (VLBI) has discovered systematic proper motions in young massive star forming regions, which in principle can be taken into account to improve kinematic distance estimates \\citep{reid09}.  A second challenge in the use of kinematic distances arises from an ambiguity between two distances (a ``near'' distance and a ``far'' distance) for sources located within the solar circle in the first and fourth Galactic quadrants. The most popular method to resolve the kinematic distance ambiguity is to measure an absorption spectrum towards the source (which has to have associated continuum radiation), and compare the velocities of the absorption lines with that of the source and the tangent point (see Fig. 1 of \\citealt{kolp03} for an illustration of this technique). Distance ambiguities have been successfully resolved towards ultracompact \\hii~regions using 21 cm \\hi~absorption \\citep{kuch94,kolp03,fish03} and 6 cm formaldehyde absorption \\citep{aray02,wats03,sewi04}.  \\hi~absorption at 21 cm has been used by \\citet{pand08} to resolve the distance ambiguity towards 34 6.7 GHz methanol masers that have either directly associated 21 cm continuum, or belong to a cluster harboring a 21 cm continuum source.  However, most methanol masers do not have any detectable radio continuum, presumably due to the young age of the exciting massive young stellar object. Hence, absorption line experiments can be used to resolve distance ambiguities towards only a small sample of 6.7 GHz methanol masers.  \\citet{burt78} and \\citet{lisz81} discovered that several \\hi~self-absorption features in \\hi~maps of the Galactic plane correlated with CO emission features, and hence hypothesized that \\hi~self-absorption could be used to determine distances to molecular clouds. \\hi~self-absorption arises from cold \\hi~in the foreground absorbing warmer radiation of the background at the a specific radial velocity within the velocity range covered by the background gas. Hence, only molecular clouds at the near kinematic distance can display \\hi~self-absorption, since at the far distance there is no background emission at the radial velocity of the cloud. The theoretical work of \\citet{flyn04} shows that molecular clouds have enough opacity in cold \\hi~to exhibit self-absorption against strong 21 cm backgrounds (such as in the Galactic plane). \\hi~self-absorption has been used by \\citet{jack02} and \\citet{busf06} to resolve the distance ambiguity towards molecular clouds and massive young stellar candidates.  In this paper, we describe work to resolve the distance ambiguity towards 6.7 GHz methanol masers detected in the AMGPS using \\hi~self-absorption. The sources were observed in CO ($J=2-1$) and NH$_3$ to determine their systemic velocities and the line profiles of thermally excited molecular emission. The maser emission of most sources have multiple emission components, which at times are spread over as much as 20~\\kms. While the central velocity of the maser emission is usually within a few \\kms~of the systemic velocity, it can at times be offset by more than 10 \\kms, and is hence not as reliable as velocities derived from thermal molecular emission. With distances determined from \\hi~self-absorption, we will then look at the distribution of sources in the Galaxy and derive the methanol maser luminosity function. This in turn will be used to compare the methanol maser population in our Galaxy with that in nearby external galaxies. ",
        "conclusions": "We used \\coii~($J=2-1$) observations and VGPS \\hi~data to resolve the kinematic distance ambiguity towards 6.7 GHz methanol masers discovered in the AMGPS. The distribution of the surface density of methanol masers as a function of Galactocentric distance agrees very well with existing estimates in the literature, although we observe the absolute numbers to be more than a factor of 3 higher. The vertical distribution of the sources has a scale height that is $\\sim 3$--5 times lower than that of the Galactic thin disk, presumably reflecting the smaller scale height of newly born massive stars. The resolution of the distance ambiguity allowed us to construct a reliable estimate of the luminosity function. Its shape does not agree with that of a power law, but has a peak around $\\sim 10^{-6}~L_\\sun$ followed by a decline towards lower luminosities. The luminosity function of 6.7~GHz methanol masers also appears to be different from that of mainline OH masers. Using the luminosity function, we derive estimates for the abundance of methanol masers in the LMC, SMC and M33 compared to the Milky Way. We find the under-abundance in M33 to be a factor of 3 higher than that in the LMC in spite of its higher metallicity. Finally, the distribution of sources between near and far distances closely follows the respective volumes sampled by the survey, thus indicating that the assumption of the near kinematic distance on a statistical basis should be avoided."
    },
    "0910/0910.1766_arXiv.txt": {
        "abstract": "{} {The condensation of diffuse gas into molecular clouds and dense cores occurs at a rate driven largely by turbulent dissipation. This process still has to be caught in action and characterized.} {We observed a mosaic of 13 fields with the IRAM-PdB interferometer (PdBI) to search for small-scale structure in the \\twCO(1-0) line emission of the turbulent and translucent environment of a low-mass dense core in the Polaris Flare.  The large size of the mosaic (1'$\\times$2') compared to the resolution (4'') is unprecedented in the study of the small-scale structure of diffuse molecular gas. } {The interferometer data uncover eight weak and elongated structures with thicknesses as small as $\\approx$ 3 mpc (600 AU) and lengths up to 70 mpc, close to the size of the mosaic.  These are not filaments because once merged with short-spacings data, the PdBI-structures appear to be the sharp edges, in space and velocity-space, of larger-scale structures. Six out of eight form quasi-parallel pairs at different velocities and different position angles. This cannot be the result of chance alignment. The velocity-shears estimated for the three pairs include the highest values ever measured in regions that do not form stars (up to 780 \\kmspc). The CO column density of the PdBI-structures is in the range $N({\\rm CO)}=10^{14}$ to $10^{15}$\\cq\\ and their \\HH\\ density, estimated in several ways, does not exceed a few 10$^3$ \\cc. Because the larger scale structures have sharp edges (with little or no overlap for those that are pairs), they have to be thin layers of CO emission. We call them SEE(D)S for Sharp-Edged  Extended (Double) Structures. These edges mark a transition, on the milliparsec scale, between a CO-rich  component and a gas undetected in the \\twCO(1-0) line because of its low CO abundance, presumably the cold neutral medium.  } {We propose that these SEE(D)S are the first directly-detected manifestations of the intermittency of interstellar turbulence. The large velocity-shears reveal an intense straining field, responsible for a local dissipation rate several orders of magnitude above average, possibly at the origin of the thin CO layers. }   \\authorrunning{Falgarone et al.} \\offprints{E.~Falgarone} \\titlerunning{Extreme velocity-shears and CO on milliparsec scale} ",
        "introduction": "\\begin{figure} \\centering \\includegraphics[width=\\hsize{}]{10963f1.eps} \\caption{The location of the 13-field mosaic observed at the Plateau de Bure interferometer (centered at RA=01:55:12.26 and Dec=87:41:56.30) is shown as the box on top of the integrated emission of the \\twCO{} and \\thCO{} \\Jone{} maps obtained at the IRAM-30m. This is a place of low, almost featureless, CO line brightness.  The arc-like structure visible in \\thCO\\ traces the outer layers of the low-mass dense core. Contour levels are shown in the wedges.} \\label{fig:location} \\end{figure}   Turbulence in the interstellar medium (ISM) remains a puzzle in spite of dedicated efforts on observational and numerical grounds. This is because it is compressible, magnetized, and multi-phase, but also because of the huge range of scales separating those of injection and dissipation of energy. Moreover, because turbulence and magnetic fields are the main support of molecular clouds against their self-gravity, turbulent dissipation is a key process among all those eventually leading to star formation \\citep[see the reviews of][] {elmescalo04,scaloelme04}.  In molecular clouds, turbulence is observed to be highly supersonic with respect to the cold gas.  It is thus anticipated to dissipate in shocks in a cloud-crossing time (\\ie\\ $\\approx$ a few 10~Myr for giant molecular clouds of 100~pc with internal velocity dispersion of a few \\kms). Magnetic fields do not significantly slow the dissipation down \\citep{mmml98}.  Actually, this is the basis of the turbulent models of star formation \\citep{maclow04} -- one of the two current scenarii of low-mass star formation -- in which self-gravitating entities form in the shock-compressed layers of supersonic turbulence.  However, while it is unquestionable that the ISM is regularly swept by large-scale shock-waves triggered by supernovae explosions that partly feed the interstellar turbulent cascade  \\citep{joung06,avillez07}, the smallest scales, barely subparsec in these simulations, are still orders of magnitude above the smallest observed structures and are unlikely to provide a proper description of the actual dissipation processes. Whether turbulent dissipation occurs primarily in compressive (curl-free) or in solenoidal (divergence-free) modes in  the interstellar medium has therefore to be considered as an open issue.  An ideal target to study turbulent dissipation is the diffuse molecular gas because it is the component in which dense cores form, with less turbulent energy density than their environment. The word ``diffuse'' here comprises all material in the neutral ISM at large that is not in dense cores \\ie\\ whose total hydrogen column density is less than a few $10^{21}$ \\cq.  This includes the mixture of cold and warm neutral medium (CNM and WNM), the edges of molecular cloud complexes (also called translucent gas), and the high latitude clouds. Diffuse gas builds up a major mass fraction of the ISM.  Actually, on the 30~pc scale, \\cite{goldsmith08} find that half the mass of the Taurus-Auriga-Perseus complex lies in regions having \\HH\\ column density below $2.1 \\times 10^{21}$ \\cq.   Turbulent dissipation may also provide clues to the ``outstanding mysteries'' raised by observations of diffuse molecular gas \\citep[see the review of][]{snow06}: the ubiquitous small scale structure, down to AU-scales \\citep{heiles07}, the remarkable molecular richness found in this hostile medium, weakly shielded from UV radiation \\citep[e.g.][]{lilu98,gredel02},  the bright emission in the \\HH\\  pure rotational lines exceeding the predictions of photon-dominated region  (PDR)  models \\citep{falgarone05,lacour05}, the \\twCO\\ small-scale structures with a broad range of temperatures, \\HH\\ densities and linewidths that preclude a single interpretation in terms of cold dense clumps \\citep{ingalls2000, ingalls2007, heithausen04, heithausen06, sakamoto03}.  The present paper extends the investigation of of turbulence down to the mpc-scale in the translucent environment of a low-mass dense core of the Polaris Flare. Over the years, this investigation has progressed along three complementary directions: \\\\ {\\it (i)} A two-point statistical analysis of the velocity field traced by the \\twCO\\ line emission, and conducted on maps of increasing size. Using numerical simulations of mildly compressible turbulence, \\cite{lis96} and~\\cite{pety03} first proposed that the non-Gaussian probability distribution functions ({\\it pdf}s) of line centroid velocity increments (CVI) be the signatures of the space-time intermittency of turbulence \\footnote{ Intermittency here refers to the empirical property of high Reynolds number turbulence to present  an excess of rare events compared to Gaussian statistics, this excess being increasingly large as velocity fluctuations at smaller and smaller scales are considered \\citep[see the review of][]{anselmet01}. Although the origin of intermittency is still an open issue \\citep[but see][]{mordant02,chevillard05,arneodo08}, it is quantitatively characterized by the anomalous scaling of the high-order structure functions of the velocity and the shape of non-Gaussian {\\it pdf}s of quantities involving velocity derivatives \\citep[e.g.][]{frisch95}.} because the extrema of CVI (E-CVI) trace extrema of the line-of-sight average of the modulus of the {\\it pos} vorticity.  Statistical analysis conducted on parsec-scale maps in two nearby molecular clouds have revealed that these extrema form parsec-scale coherent structures~\\citep[resp. Paper~III, HF09]{hily08,hf09}. \\\\ {\\it (ii)} A detailed analysis (density, temperature, molecular abundances) of these coherent structures, based on their molecular line emission. The gas there is more optically thin in the \\twCO\\ lines, warmer and more dilute than the bulk of the gas \\citep[hereafter Paper II]{hily07}, and  large \\HCOp\\ abundances, unexpected in an environment weakly shielded from UV radiation, have been detected there \\citep[Paper~I]{falgarone06}.\\\\ {\\it (iii)} Chemical models of non-equilibrium warm chemistry triggered by bursts of turbulent dissipation~\\citep{joulain98}. The most recent progresses along those lines include the chemical models of turbulent dissipation regions (TDRs) by \\cite{godard09} and their successful comparison to several data sets, among which new submillimeter detections of \\thCHp(1-0) (Falgarone et al. in preparation).  The \\twCO\\Jtwo\\ observations of the Polaris Flare with unprecedented angular resolution and dynamic range are the first to evidence the association between extrema of CVI and observed velocity-shears\\footnote{We use velocity-shear rather than velocity-gradient because the observations provide cross-derivatives of the velocity field, \\ie\\ the displacement measured in the plane-of-the-sky ({\\it pos}) is perpendicular to the line-of-sight velocity} (HF09). No shock signature (density and/or temperature enhancement, SiO detection) has been found in the coherent structure of E-CVI identified in the Polaris Flare (Hily-Blant and Falgarone, in preparation). All the above suggest (but does not prove yet) that the coherent structures carrying the statistical properties of intermittency are regions of intense velocity-shears where dissipation of turbulence is concentrated.  The \\twCO(1-0) observations reported in this paper have been performed in a field located on one branch of the Polaris Flare E-CVI structure, in the translucent and featureless environment of a dense core (Fig.~\\ref{fig:location}).  The outline of the paper is the following: the observations and data reduction are described in Section 2. The observational results are given in Section 3. The characterization of the emitting gas is made in Section 4 and we discuss, in Section 5, the possible origin and nature of the CO structures that we have discovered. Section 6 puts our results in the broad perspective provided by other data sets and Section 7 compares them to  chemical model predictions and numerical simulations of turbulence. The conclusions are given in Section 8. ",
        "conclusions": "IRAM-PdBI observations of a mosaic of 13 fields in the turbulent environment of a low-mass dense core have disclosed small and weak \\twCO(1-0) structures in translucent molecular gas. They are straight and elongated structures but they are not filaments because, once merged with short-spacings data, the PdBI- structures appear as the sharp edges of larger-scale structures. Their thickness is as small as $\\approx$ 3 mpc (600 AU), and their length, up to 70 mpc, is only limited by the size of the mosaic. Their CO column density is a well determined quantity for the excitation conditions found at larger scale and is in the range $N({\\rm CO)}=10^{14}$ to $10^{15}$\\cq.  Their \\HH\\ density, estimated in several ways, including the continuum emission of the brightest structure, does not exceed a few 10$^3$~\\cc. Their well-distributed orientations can be followed in the larger-scale environnement of the field. Six of them form three pairs of quasi-parallel structures, physically related. The velocity-shears estimated for the three pairs include the largest ever measured in non-star-forming clouds (up to 780 \\kmspc).  The PdBI-structures are therefore not isolated and are the edges of so-called SEE(D)S for sharp-edged extended (double) structures.  We show that the SEE(D)S are thin layers of CO-rich gas and that their sharp edges pinpoint a  small-scale dynamical process, at the origin of the CO contrast detected by the PdBI.  We propose that the SEE(D)S are the outcomes of the chemical enrichment driven by intense  dissipation occurring in large velocity-shears and that they are CO-rich layers swept along by the straining field of CNM turbulence.   The present work is the first detection of mpc-scale intense velocity-shears belonging to a parsec-scale shear. The large departure from average of the kinematic properties of these structures, confirms that they are a manifestation of the small-scale intermittency of turbulence in this high latitude field, a property  already established on statistical grounds (HF09). The values of the velocity-shears (or rate-of-strain) provide a quantitative constraint on the dissipation rate that can be  compared to chemical models. The link between the turbulent dissipation in the diffuse gas and the dense core observed in the vicinity of the PdBI mosaic (Fig.~\\ref{fig:location}) still remains to be established.  Last, we would like to stress that sub-structure still exists in these mpc-scale structures of the diffuse ISM and that the next generation of interferometers (e.g. ALMA) should be able to observe gas  at the dissipation scale of turbulence (that is still unknown) or at least observe the effects on the ISM (temperature, excitation, molecular abundances) of the huge release of energy expected to occur there."
    },
    "0910/0910.1882_arXiv.txt": {
        "abstract": "We analyze the time dependence of fluid variables in general relativistic, magnetohydrodynamic simulations of accretion flows onto a black hole with dimensionless spin parameter $a/M=0.9$. We consider both the case where the angular momentum of the accretion material is aligned with the black hole spin axis (an untilted flow) and where it is misaligned by $15^\\circ$ (a tilted flow). In comparison to the untilted simulation, the tilted simulation exhibits a clear excess of inertial variability, that is, variability at frequencies below the local radial epicyclic frequency. We further study the radial structure of this inertial-like power by focusing on a radially extended band at $118(M/10M_\\odot)^{-1}{\\rm Hz}$ found in each of the three analyzed fluid variables. The three dimensional density structure at this frequency suggests that the power is a composite oscillation whose dominant components are an over dense clump corotating with the background flow, a low order inertial wave, and a low order inertial-acoustic wave. Our results provide preliminary confirmation of earlier suggestions that disk tilt can be an important excitation mechanism for inertial waves. ",
        "introduction": "Quasi-periodic oscillations (QPOs) are observed in the X-ray light curves of many black hole X-ray binaries (see \\citealt{rem06} for a recent review).  They have also been observed in extragalactic ultraluminous X-ray sources \\citep{str03,str07} and in one active galactic nucleus \\citep{gie08}.  The origin of these phenomena is still far from clear.  One class of models centers on trapped wave modes within the accretion flow.  In particular, a variety of modes have been proposed in hydrodynamic models of geometrically thin accretion disks  (e.g. \\citealt{wag99,kat01}). Axisymmetric inertial-acoustic modes (``$f$-'' or ``inner $p$-modes'') and axisymmetric inertial modes (``$g$-'' or ``$r$-modes'') \\footnote{Throughout this paper we will use the term inertial mode and $r$-mode interchangeably.  We prefer to avoid the use of the term $g$-mode, whose primary restoring force is generally due to entropy gradients.  In disks, the primary restoring force for inertial modes is due to specific angular momentum gradients.} can be trapped due to the existence of the innermost stable circular orbit around a black hole, which produces a maximum in the radial profile of the radial epicyclic frequency.  Non-axisymmetric inertial modes can also be radially trapped between inner and outer Lindblad resonances.  The very existence of inertial modes in accretion disks has been challenged recently, as numerical simulations appear to show that magnetorotational (MRI) turbulence suppresses them.  This has been demonstrated in local shearing box simulations \\citep{arr06}, whose boundary conditions would artificially trap a discrete axisymmetric mode spectrum.  Power spectra from these simulations show discrete acoustic modes and a radial epicyclic oscillation, but no inertial modes. This behavior has also been seen in global simulations of accretion disks in a pseudo-Newtonian potential: hydrodynamic disks exhibit trapped axisymmetric inertial modes, whereas MHD turbulent disks do not \\citep{rey08}. It appears that axisymmetric inertial modes, which necessarily have frequencies at or below the local radial epicyclic frequency, are particularly vulnerable to nonlinear damping by MRI turbulence, which has a power spectrum that peaks near the orbital frequency.  Non-axisymmetric inertial modes can have higher frequencies and might be less vulnerable to MRI turbulence, but there has as yet been no convincing demonstration of the existence of a trapped inertial mode spectrum in any MHD simulation.  Another reason why discrete inertial modes may not be present in accretion disks with MRI turbulence is that even subthermal magnetic fields can extend the inner trapping radius of the propagation zone of both axisymmetric and non-axisymmetric inertial modes down to the innermost stable circular orbit (ISCO) \\citep{fu09}.  Unless inertial waves can reflect off the plunging region of the flow, standing waves will no longer be sustainable.  On the other hand, it has been suggested that warps and eccentricity in disks may play a fundamental role in a nonlinear excitation mechanism of inertial modes in geometrically thin accretion disks \\citep{kat04a,kat08,fer08}. These large scale deformations interact with trapped $r$-modes giving rise to intermediate modes. These intermediate modes then couple back to the warp or eccentricity to produce positive feedback on the original $r$-mode oscillations (see \\citealt{fer08} for a more detailed explanation of this coupling mechanism).  \\cite{fra07} recently completed fully general relativistic, MHD simulations of accretion disks with misaligned black hole spin and disk angular momentum vectors (``tilted disks''). These tilted disks are strongly warped near the black hole and exhibit global epicyclic oscillations superimposed on the turbulence in the flow \\citep{fra08}. These oscillations manifest themselves as eccentric orbits of fluid elements in the disk, albeit with a $180^\\circ$ flip in orientation of the elliptical orbits across the midplane of the disk. Tilted disks may therefore provide favorable conditions for the nonlinear excitation of inertial modes.  Here we report the results of a search for trapped modes in simulations of both an untilted and tilted disk. In agreement with previous work, we do not observe the presence of modes in the untilted simulation. However, at the same numerical resolution, we observe significant excess power with frequencies characteristic of inertial waves in the tilted disk. Because of the complexity of the background flow, the physical nature of this power is difficult to determine. Nevertheless, its spatial structure appears to confirm that this is partly inertial in character. This may represent preliminary confirmation of the warp/eccentricity excitation mechanism, showing that it can be strong enough to overcome damping due to MRI turbulence. It also shows that trapping, in the sense of radial localization of power, can be maintained even in the presence of magnetic fields.  This paper is organized in the following manner. In section 2 we briefly review the simulation parameters and discuss our power spectrum analysis of the simulation data. Section 3 presents power spectra from both untilted and tilted geometries, evidence that the latter may be dominated by inertial-like variability, and an analysis of the three dimensional structure of the power at a particular frequency within this inertial regime. We summarize our conclusions in section 4. ",
        "conclusions": "Previous work has demonstrated that inertial modes are vulnerable to disruption by MRI turbulence \\citep{arr06,rey08}. Our Fourier analysis of the untilted simulation, a configuration devoid of tilt or eccentricity, corroborates this conclusion. However, nonlinear analytic calculations \\citep{kat04a, kat08, fer08} have indicated that disk warps and eccentricity may excite inertial modes; both of these features are generically present in tilted accretion flows. Though the analysis of the variability in our tilted simulation at $118{\\rm Hz}$ does not appear to be wholly inertial (or even acoustic) in nature, the presence of two weak, odd parity, $m=0$ and $m=2$ components provides some support for this hypothesis.  Evidence that disk tilt is connected with the excitation of inertial or acoustic waves is twofold. First, fundamentally different shapes in the spectra from the untilted and tilted simulations implies a correlation between tilt and the power at frequencies characteristic of inertial waves. Second, the radial structure of the variability suggests a superposition of modes, at least one of which may be described as inertial in nature in the context of relevant analytic models. Together, these clues establish the presence of variability unique to a tilted geometry which appears to be at least partially inertial in character.  Recalling that a tilted disk configuration naturally yields a low frequency QPO activity in black hole X-ray binaries and speculating that our work's inertial-like variability may be related to high frequency QPO activity in the same systems, we advocate further numerical exploration of these tilted geometries. Torques associated with the Kerr spacetime can produce a bodily precession of the tilted disk at frequencies within the $0.1$ and $30{\\rm Hz}$ range associated with the low frequency QPO. Because this precession frequency is strongly dependent on the radial extent of the disk in this model, it readily explains the observed correlation between the LFQPO's frequency and disk flux \\citep{rem06,fra08,ing09}. Similarly, our current work shows inertial or acoustic variability excited and maintained by the disk tilt that exhibits frequencies of order the radial epicyclic frequency maximum. Such frequencies fall within the observed range of $40$ to $450{\\rm Hz}$ for HFQPOs in stellar mass black hole systems.   HFQPOs are only observed when a system's X-ray emission exhibits the characteristics of the steep power-law state (cf. \\citealt{rem06} for a review of the properties of X-ray black hole binaries), and we must acknowledge, however, that our simulations are almost certainly not accurate representations of that state.  For instance, since they neither account for radiative cooling nor capture all forms of dissipation, our simulations likely do not reproduce some of the key features found in various models for the state's structure (e.g. the energetically significant corona in \\citealt{don07}). However, we are confident that our simulations and analysis provide insight into the dynamical effects of a tilted accretion flow, particularly, the excitation of inertial and acoustic modes.  \\sloppypar{ We thank Chris Done, Paul Henisey, Mami Machida, Phil Marshall, Gordon Ogilvie, Chris Reynolds, and Alexander Tchekhovskoy for useful conversations, and Sam Cook for his help with data transport and preparation. We also thank the anonymous referee for comments that significantly improved this paper. This work was supported in part by NSF grant AST-0707624.  The work of BTF was supported by FCT (Portugal) through grant SFRH/BD/22251/2005. PCF acknowledges support from NSF grant AST-0807385 and an American Astronomical Society International Travel Grant. We are also grateful to the Nordic Institute for Theoretical Astrophysics for hosting a workshop on QPOs where much of this work was completed. }"
    },
    "0910/0910.1599_arXiv.txt": {
        "abstract": "The joint likelihood of observable cluster signals reflects the astrophysical evolution of the coupled baryonic and dark matter components in massive halos, and its knowledge will enhance cosmological parameter constraints in the coming era of large, multi-wavelength cluster surveys.  We present a  computational study of intrinsic covariance in cluster properties using halo populations derived from Millennium Gas Simulations (MGS).  The MGS are re-simulations of the original $500 \\hinv \\mpc$ Millennium Simulation performed with gas dynamics under two different physical treatments: shock heating driven by gravity only ($\\GO$) and a second treatment with cooling and preheating ($\\PH$).   We examine relationships among structural properties and observable X-ray and Sunyaev-Zel'dovich (SZ) signals for samples of thousands of halos with $M_{200} \\ge 5 \\times 10^{13} \\hinv\\msun$ and $z < 2$.     While the X-ray scaling behavior of \\PH model halos at low-redshift offers a good match to local clusters, the model exhibits non-standard features testable with larger surveys, including weakly running slopes in hot gas observable--mass relations and $\\sim 10\\%$ departures from self-similar redshift evolution for $10^{14} \\hinv\\msun$ halos at redshift $z \\sim 1$.    We find that the form of the joint likelihood of signal pairs is generally well-described by a multivariate, log-normal distribution, especially in the \\PH case which exhibits less halo substructure than the \\GO model.  At fixed mass and epoch, joint deviations of signal pairs display mainly positive correlations, especially the thermal SZ effect paired with either hot gas fraction ($r=0.88/0.69$ for $\\PH/\\GO$ at $z=0$) or X-ray temperature ($r=0.62/0.83$).    The levels of variance in X-ray luminosity, temperature and gas mass fraction are sensitive to the physical treatment, but offsetting shifts in the latter two measures maintain a fixed $12\\%$ scatter in the integrated SZ signal under both gas treatments.  We discuss halo mass selection by signal pairs, and find a minimum mass scatter of $4\\%$ in the \\PH model by combining thermal SZ and gas fraction measurements. ",
        "introduction": "Accurate cosmology using surveys of clusters of galaxies requires a robust description of the relations between observed cluster signals and underlying halo mass.  Even without strong prior knowledge of the mass-signal relation, cluster counts, in combination with other probes, add useful constraining power to cosmological parameters \\citep{cunha:09a}.   However, significant improvements can be realized when the error in mass variance is known \\citep{limaHu:05, cunha:09b}.   Improvements can also be gained  by extending the model to multiple observed signals \\citep{cunha:08}, especially when an underlying physical model can effectively reduce the dimensionality of the parameter sub-space associated with the model \\citep{younger:06}.   The coming era of multiple observable signals from combined surveys in optical, sub-mm and X-ray wavebands invites a more holistic approach to modeling multi-wavelength signatures of clusters.  Signal covariance characterizes survey selection, in terms of mass and additional observables. For the case of X-ray selected samples, \\cite{nord:08} demonstrate that luminosity--temperature covariance can mimic apparent evolution in the luminosity--mass relation under analysis that combines deep, X-ray-flux limited samples with local, shallow ones.  Employing a selection observable with small mass variance minimizes such errors.  Recent work has shown that the total gas thermal energy, $Y$, observable via an integrated Sunyaev-Zeldovich (SZ) effect \\citep{carlstrom:04} or via X-ray imaging and spectroscopy, is a signal that scales as a power-law in mass with only $\\sim 15\\%$ scatter \\citep{white:02, kravtsov:06, maughan:07, ohara:07, zhang:08, jeltema:08}.   However, unbiased estimates of the mass selection function for $Y$ or any other signal requires accurate knowledge of how the signal--mass scaling relation evolves with redshift.   The redshift behavior of signals is generally not well known empirically, although recent work has begun to probe evolution in X-ray signals to $z \\sim 1$ \\citep{maughan:06,vikhlinin:08}.  Emerging samples from wide-area SZ surveys should dramatically improve this situation.  One can address signal--mass covariance using hydrostatic, virial or lensing mass estimates from observations, but several sources of systematic and statistical error challenge this approach.   For hydrostatic masses, early gas dynamic simulations \\citep{evrard:90,nfw, thomas:97} suggested that turbulent gas motions drove hydrostatic masses to underestimate true values by $\\sim 20\\%$.  More sophisticated recent models, with a factor thousand improvement in mass resolution,  demonstrate this effect at a similar level in the mean, with $\\sim 15\\%$ scatter among individual systems \\citep{rasia:06, nagai:06, jeltema:08}.  Cluster masses can also be measured by the shear induced on background galaxies due to gravitational lensing.  With this method, individual  cluster masses have mass uncertainties $\\sim 20\\%$ due to cosmic web confusion \\citep{hoekstra:03, deputter:05}, but large samples can reduce the uncertainty in the mean.  The mean scaling behavior of samples binned in some selection signal offers another empirical path to measuring covariance. Non-zero covariance between the selection and an independent, follow-up signal implies that the selection-binned scaling relation of the followup signal with mass need not match that signal's intrinsic mass scaling.  Comparison of scaling relations from differently selected samples thereby offers insight into covariance.   \\cite{rykoff:08} offer a first attempt at this exercise for X-ray luminosity and optical richness using the optically-selected SDSS \\maxbcg\\  sample \\citep{koester:07}.  The sample contains $\\sim 13,000$ clusters for which weak lensing mass estimates have been made by stacking the shear of richness-binned sub-samples \\citep{sheldon:07,johnston:07}.   \\cite{rykoff:08} stack Rosat All-Sky Survey data \\citep{voges:99} in the same \\maxbcg\\  richness bins, and find that the mean X-ray luminosity--mass relation derived with richness binning is consistent at the $\\sim 2 \\sigma$ level with relations derived solely from X--ray data \\citep{reiprich:02, stanek:06}.  A theoretical approach to studying cluster covariance is to realize populations via numerical simulation.  While high resolution treatment of astrophysical processes, including star formation, supernova and AGN feedback, galactic winds and thermal conduction have been included in recent simulations \\citep{dolag:05, borgani:06, kravtsov:06, sijacki:08, puchwein:08}, the computational expense has limited sample sizes to typically a few dozen objects.   A detailed study of population covariance requires larger sample sizes, as can be generated by lower resolution simulations of large cosmological volumes using a more limited physics treatment \\citep{bryan:98, hallman:06, gottlober:07}.  We take the latter approach in this paper, focusing on the bulk properties of massive halos identified the Millennium Gas Simulations (MGS), a pair of resimulations of the original $500 \\hinv$ Mpc Millennium run \\citep{springel:05}, each with $10^9$ total particles, half representing gas and half dark matter.  The pair of runs use different treatments for the gas physics --- a gravity-only ({\\texttt{GO}}) simulation, sometimes called ``adiabatic'',  in which entropy is increased via shocks, and a simulation with cooling and preheating  ({\\texttt{PH}}).  The former ignores galaxies as both a sink for baryons and a source of feedback for the hot intracluster medium (ICM).  The latter also ignores the mass fraction contribution of galaxies, but it approximates the feedback effects of galaxy formation by a single parameter, an entropy level imposed as a floor at high redshift  \\citep{evrard:91, kaiser:91, bialek:01, kay:07, gottlober:07}.  Our study focuses on samples of $\\sim 5000$ halos with mass $M > 5 \\times 10^{13} \\hinv \\msun$ examined at multiple epochs covering the redshift range $0 < z < 2$.  The paper is organized as follows:  in Section \\ref{sec:sims} we discuss the details of the Millennium Gas Simulations and our halo finding approach. In Section \\ref{sec:scaling}, we present mean scaling relation behavior in both mass and redshift.  We then turn to covariance about the mean scaling relations in Section \\ref{sec:covar}.  Unless otherwise noted, our units of mass are $10^{14} \\hinv \\msun$ and halo mass is defined within a sphere encompassing a density contrast $\\Delta_c = 200$ times the critical density. ",
        "conclusions": "\\label{sec:con}  We analyze scaling relations for multiple properties of massive halos taken from a pair of gas dynamic simulations with different gas physics treatments.   Our samples contain tens of thousands of halos with masses $M_{200} > 5 \\times 10^{13} \\hinv \\msol$ at redshifts $z \\lta 2$.  The physical treatments of gravity only ($\\GO$) or preheating ($\\PH$) are both highly idealized, but we show that the latter reproduces the scaling relation behavior of core-extracted X-ray measures of local clusters.  The dark matter velocity dispersion scales with mass and redshift according to self-similar expectations, indicating that the virial theorem is respected regardless of gas treatment.  However, the gas behavior in both treatments differs from self-similarity.   The deviations in the \\GO case tend to be small and are related to the mass-limited sample definition.  The entropy injection of the \\PH model drives more substantial deviations from self-similar scaling.  At $z=0$, the ICM mass fraction varies with mass in a manner roughly consistent with observed measurements of local clusters.   We find that $\\ficm$ requires a quadratic fit in log-mass, and mild curvature in the logarithmic scalings of $Y$ and $\\elbol$ as a function of mass is also evident.  While the ICM in \\PH halos is lower in mass compared to the \\GO case, it is also slightly hotter.   The effects on $\\ficm$ and $T_m$ nearly cancel when combined to form the thermal SZ signal, leading to nearly identical $Y$--$M$ scaling relations above $3 \\times 10^{14}$ for both \\PH and \\GO cases at low redshift.  The ICM mass fraction at fixed mass declines weakly with redshift, by $10\\%$ in $5 \\times 10^{14} \\hinv\\msol$ halos at $z=1$ and with larger reductions at lower masses.  While {\\sl Chandra\\/} observations of optically-selected clusters in the RCS survey show evidence of reduced gas fractions at $z=1$ \\citep{hicks:06}, further work is needed to address the quantitive level of agreement between the observations and \\PH model expectations. The \\PH baryon fraction evolution drives departures from self-similar predictions in the  $Y$--$M$ and $\\elbol$--$M$ relations;  slopes tend to steepen slightly and the normalization at $10^{14} \\hinv\\msol$ is lower than self-similar expectations at high redshift.  We present the first systematic investigation of property covariance in the computational samples of massive halos.  The data generally support a multivariate log-normal form for the joint distribution of signals at fixed mass and redshift.   All measures depart somewhat from  an exact gaussian form, but the deviations in gas measures are smaller for the \\PH model due to the suppression of substructure caused by the preheating.   Most signal pairs exhibit positive correlations, with the lone exception of $-0.1$ between $\\sigmadm$  and $\\ficm$ in the \\PH case.  The thermal SZ signal displays a robust $13\\%$ scatter that is strongly correlated with variations in both ICM gas mass and temperature, with $\\ficm$ dominating in the \\PH case and $T_m$ being more important in the \\GO treatment.  Combining multiwavelength observations offers an opportunity to improve selection of clusters by their intrinsic mass.   We derive the mass variance of signal pairs and show that combining strongly correlated signals always improves mass selection, even when one of the signals by itself is a comparatively poor mass proxy.   The combination of thermal SZ and ICM mass fraction in the \\PH case selects halo mass with just $4\\%$ intrinsic rms scatter.  Identifying the root causes behind the terms in the covariance matrix is a considerable task that we leave for future work.   Mergers \\citep{roettigerBurnsLoken:97, rickerSarazin:01}  and assembly bias (\\cite{Boylan-Kolchin:09} and references therein) will surely play important roles for many cluster signals.  Studies of merger history behavior will shed light on survey selection properties, particularly potential biases related to the dynamical state of a halo.  The simple treatment of baryon physics in our simulations limits our investigation to the hot, thermal ICM of clusters.  More extensive physical treatments that incorporate galaxy and supermassive black hole formation and other physics such as MHD and non-thermal plasmas will ultimately extend the set of observable halo signals into the optical/NIR and radio wavebands."
    },
    "0910/0910.1550_arXiv.txt": {
        "abstract": "In the weak field limit general relativity reduces, as is well known, to the Newtonian gravitation. Alternative theories of gravity, however, do not necessarily reduce to Newtonian gravitation; some of them, for example, reduce to Yukawa-like potentials instead of the Newtonian potential. Since the Newtonian gravitation is largely used to model with success the structures of the universe, such as for example galaxies and clusters of galaxies, a way to probe and constrain alternative theories, in the weak field limit, is to apply them to model the structures of the universe. In the present study we consider how to probe Yukawa-like potentials using $N$-body numerical simulations.  \\PACS{04.50.+h, 04.50.Kd, 45.50.-j} ",
        "introduction": "\\label{intro} Einstein's General Relativity (GR) is one of the most beautiful theories ever imagined by the human mind. Although GR is a successful theory of gravitation, it is unable to explain, for example, the accelerating expansion of the universe, unless a cosmological constant or a dark energy fluid is considered. Nowadays other theories of gravitation intend to give an alternative interpretation to this accelerating expansion (see, e.g., \\cite{piazzamarinoni}). Most of these alternative theories, although having different approaches (some scalar-tensor theories of gravitation, nonsymmetric gravitational theory, etc), reduce, in the weak-field limit, to a Yukawa-like gravitational potential (hereafter YGP), i.e., \\begin{equation} \\phi(r)=-\\frac{GM}{r}e^{-r/{\\lambda}}. \\label{yukpot} \\end{equation}  The above equation gives the potential of a point mass m at a distance r; G is the universal gravitational constant and $\\lambda$ is a constant. When $\\lambda \\rightarrow \\infty$, this potential tends to the Newtonian one. Note that the parameter $\\lambda$ is the Compton wavelength of the exchange particle, which in the present case is a graviton. The graviton mass is related to  $\\lambda$ through the well known equation $m_g=h / \\lambda c$, where $h$ is the Planck constant and c is the speed of light.  The YGP has been investigated in the literature, in particular, in galactic astronomy and cosmology. For example, Signore \\cite{signore} studies this potential under the cosmological context, on the other hand, de Araujo \\& Miranda \\cite{araujoemiranda2007} study how variations of the $\\lambda$ parameter can disturb galactic disks. This last study generalizes the (Newtonian) Toomre's \\cite{toomre} expressions for rotation curves in order to account for a YGP. Rodriguez-Meza et al. \\cite{rodriguezmeza} consider the YGP as a correction on the Newtonian potential.  Recently, some numerical approaches have been developed to study alternative theories. Cervantes-Cota et al. \\cite{cervantes2007a,cervantes2007b} have developed numerical studies of a scalar-tensor theory in the weak field limit, using the $N$-body method. In these studies they consider that the gravitational potential in the weak field limit is  given by  $$\\Phi_{STT} = \\phi_N(r) +  \\alpha \\phi(r)/ (1+ \\alpha)\\, , $$  \\par\\noindent where $\\Phi_{STT}$ is the potential from the scalar-tensor theory, $\\phi_N(r)$ is the Newtonian potential, $\\phi(r)$ is given by Eq.(1), and $\\alpha \\equiv 1/(3 + 2 \\omega)$, with $\\omega$ being the Brans-Dicke parameter \\cite{brans}. They changed a  particular $N$-body code replacing the Newtonian potential by the potential $\\Phi_{STT}$ and analyzed its influence on isolated galaxies, two colliding spiral galaxies and issues concerning the formation of bars. Simulations of cosmological structure formation in $\\rm{\\Lambda}$CDM scenarios are also studied.  In the present paper, we use similar techniques to that used by Cervantes et al. \\cite{cervantes2007a,cervantes2007b}, although with a gravitational potential derived from other theories. We modify a popular $N$-body code, \\textbf{Gadget-2} \\cite{springel2005}, \\textsl{replacing} the Newtonian potential by a pure YGP as given by Equation \\ref{yukpot}, which is derived from some alternative theories of gravitation (see, e.g., Visser's theory \\cite{visser}).  It is worth stressing that although we have chosen a particular gravitational potential, the approach considered here can be applied to any alternative gravitational potential.  This paper is organized as follows: in section 2, we show the changes made in the \\textbf{Gadget-2} code, in section 3, we test the YGP \\textbf{Gadget-2} code using a pair of particles and $N$-body systems. Finally, in section 4, we present the conclusions and briefly discuss how this modified code can be used to test alternative theories of gravitation via galactic dynamics. ",
        "conclusions": "The present paper considers the use of the well known \\textbf{Gadget-2}, with the appropriate modifications, to study alternative theories of gravitation, with particular attention to the YGP. In this first paper we discuss how to modify the \\textbf{Gadget-2} and evaluate how reliable this modified code is, performing a series of tests.  As we have shown in the previous sections, our tests are successful: numerical errors from our changes on \\textbf{Gadget-2} are neglected and we can use this modified code to probe alternative theories of gravitation. For example, we can use arguments based on galactic dynamics to exclude or not the Yukawa potential hypothesis. This is very important, due to the fact that alternative theories appear frequently in the literature, but so far there are no many investigations in  the galactic scales using $N$-body codes and galactic dynamics formalism. We will show, in future papers, how to test these alternative theories using early and late-type galaxies."
    },
    "0910/0910.3887_arXiv.txt": {
        "abstract": "We investigate whether a tachyonic scalar field, encompassing both dark energy and dark matter-like features will drive our universe towards a Big Brake singularity or a de Sitter expansion. In doing this it is crucial to establish the parameter domain of the model, which is compatible with type Ia supernovae data. We find the 1$\\sigma$ contours and evolve the tachyonic sytem into the future. We conclude, that both future evolutions are allowed by observations, Big Brake becoming increasingly likely with the increase of the positive model parameter $k$. ",
        "introduction": "With the discovery of cosmic acceleration \\cite{cosm} the quest for modeling dark energy \\cite{dark} has started. Besides the most simple cosmological constant, other models based on various perfect fluids with negative pressure, like Chaplygin gas \\cite{Chaplygin}, minimally and non-minimally coupled scalar fields and fields having non-standard kinetic terms \\cite% {kinetic,tachyons} were advanced. The latter ones include as a subclass the models based on different forms of the Born-Infeld-type action, which is often associated with the tachyons arising in the context of string theory \\cite{string}. Due to the non-linearity of the dependence of the tachyon Lagrangians on the kinetic term of the tachyon field, the dynamics of the corresponding cosmological models appears to be very rich.  The tachyon model studied in paper \\cite{we-tach} contains a 2-fluid analogue scalar field $T$, the dynamics of which is given by a simple potential, depending on two parameters, $\\Lambda $ and $k$. The model is homogeneous and isotropic. A phase space diagram in the tachyonic field and its derivative $s\\equiv \\dot{T}$ shows 5 type of distinct cosmological evolutions possibly occurring for the model, some of them containing regimes where $s$ is superluminal. All evolutions originate from one of the Big Bangs of the model, but they either end in a de Sitter infinite exponential expansion, as the $\\Lambda $CDM model does, or in a future singularity characterized by a regular scale factor $a$, vanishing Hubble parameter $H$ and energy density $\\varepsilon $, but infinite $s$ and pressure $p$. Most notably, the second time derivative of the scale factor goes to $-\\infty $, the reason why we call this singularity a Big Brake.  A kinematical analysis \\cite{Barrow} predicted the existence of such singularities, named sudden future singularities. From a combined kinematical and observational reasoning alone, sudden future singularities could occur as early as in ten million years \\cite{Dabrowski}, however no underlying dynamics is known to support this.  Classically the Big Brake singularity is stable. This can be seen by a series expansion of the scale factor in the vicinity of the singularity and checking the stability conditions advanced in Ref. \\cite{Barrow-priv}. Its quantum study indicated singularity avoidance \\cite{KKS}.  Recently \\cite{tachyon-prd} the compatibility of the model with type Ia supernovae observation has been investigated. After we present some basic features of the model in Section 2, in Section 3 we give more details on this compatibility check, in terms of the original variables employed in Ref. \\cite{we-tach}. Then in Section 4 we stress the crucial difference between negative and positive values of the model parameter $k$. While for the former all evolutions end in the de Sitter attractor, for positive $k$ the 1$\\sigma $ contour compatible with type Ia supernovae contains both states which evolve into de Sitter or into a Big Brake. In this dynamical model the Big Brake can occur no earlier than $10^{8}$ years.  The Big Brake singularity belongs to the class of soft cosmological singularities, which also includes other representants \\cite{soft}. Other types of singularities arising in the study of various dark energy models include the Big Rip singularity \\cite{Rip}, present in some models with phantom dark energy \\cite{phantom}. The possibility of existence of a phase of contraction of the universe, ending up in the standard Big Crunch cosmological singularity was also considered \\cite{Crunch}.  \\textit{Unit convention: }the Newtonian constant is normalized as $8\\pi G/3=1 $ and we take $c=1$. ",
        "conclusions": ""
    },
    "0910/0910.4233_arXiv.txt": {
        "abstract": "We revisit the statistical significance of the ``dark flow'' presented in \\cite{kashlinsky09}.  We do not find a statistically significant detection of a bulk flow.  Instead we find that CMB correlations between the 8 WMAP channels used in this analysis decrease the inferred significance of the detection to 0.7$\\sigma$. ",
        "introduction": "\\label{sec:intro}  A recent set of papers \\citep{kashlinsky08, kashlinsky09} claims to have detected the velocities of galaxy clusters with respect to the cosmic microwave background (CMB) frame by means of the kinetic Sunyaev-Zel'dovich (kSZ) effect \\citep{sunyaev72}.  The papers suggest the existence of a ``dark flow'': a 700 km s$^{-1}$ bulk flow of all matter out to a redshift of at least $z\\simeq 0.1$ ($\\rm{r}\\simeq400$ Mpc).  The magnitude and direction of the flow are claimed to be consistent with the peculiar velocity of the Local Group with respect to the CMB frame as inferred from the CMB dipole \\citep{kogut93}.  Velocity coherence over such large scales is not predicted by the standard $\\Lambda\\rm{CDM}$ cosmology and would, if confirmed, constitute a major observational result.  In this paper we revisit the analysis presented in \\cite{kashlinsky09}, hereafter referred to as K09.  The K09 analysis seeks to measure the kSZ signal of a sample of $\\sim$700 X-ray-selected galaxy clusters.  The 3-year WMAP temperature maps for 8 differencing assemblies \\citep{hinshaw07} are high-pass filtered in an attempt to remove the primary CMB anisotropy.  The temperatures of the filtered maps at the galaxy cluster locations are fit to a dipole, which is interpreted as the kSZ signal induced by a bulk flow of the galaxy clusters.  We will argue that the uncertainty of this measurement is dominated by primary CMB anisotropy, not detector noise.  As the CMB is observed by all 8 WMAP channels, the errors are highly correlated between these channels, and the inferred detection significance is greatly reduced. ",
        "conclusions": "\\label{sec:conclusion}  We have revisited the analysis presented in \\cite{kashlinsky08, kashlinsky09} which reports a significant detection of a bulk flow of $\\sim$700 galaxy clusters out to $z\\simeq0.1$ by means of the kSZ effect.  We have demonstrated that the estimates for the kSZ signal are highly correlated between the different WMAP channels used in this analysis and that this correlation is caused by primary CMB anisotropy.  We have simulated the errors on the kSZ measurement while taking into account these CMB correlations and find that there is not a significant detection of a kSZ signal or bulk flow.   \\begin{figure*} \\begin{center} \\includegraphics[width=1.0\\textwidth]{err.pdf} \\caption{1000 simulated estimates for the 3 dipole components.  These simulations take into account the CMB correlations between the different WMAP channels.  The uncertainty is highest on the dipole components that lie in the galactic plane ($a_x$ and $a_y$) because of the geometry of the galactic mask.} \\label{fig:err} \\end{center} \\end{figure*}   \\begin{figure*} \\begin{center} \\includegraphics[width=0.7\\textwidth]{corr.pdf} \\caption{1000 simulated estimates for the $a_x$ dipole component from two randomly chosen WMAP channels, Q1 and W2.  The estimates are highly correlated ($\\rho=0.9$).  This high level of correlation is common to all pairs of channels and is caused by primary CMB fluctuations.} \\label{fig:corr} \\end{center} \\end{figure*}"
    },
    "0910/0910.3139_arXiv.txt": {
        "abstract": "A few star clusters in the Magellanic Clouds exhibit composite structures in the red-clump region of their colour--magnitude diagrams. The most striking case is NGC~419 in the Small Magellanic Cloud (SMC), where the red clump is composed of a main blob as well as a distinct secondary feature. This structure is demonstrated to be real and corresponds to the simultaneous presence of stars which passed through electron degeneracy after central-hydrogen exhaustion and those that did not. This rare occurrence in a single cluster allows us to set stringent constraints on its age and on the efficiency of convective-core overshooting during main-sequence evolution. We present a more detailed analysis of NGC~419, together with a first look at other populous Large Magellanic Cloud clusters which are apparently in the same phase: NGC~1751, NGC~1783, NGC~1806, NGC~1846, NGC~1852 and NGC~1917. We also compare these Magellanic Cloud cases with their Galactic counterparts, NGC~752 and NGC~7789. We emphasise the extraordinary potential of these clusters as {\\em absolute} calibration marks on the age scale of stellar populations. ",
        "introduction": "\\begin{figure}[b] \\begin{center} \\includegraphics[width=0.8\\textwidth]{figure_lr.eps} \\caption{False-colour images (in the electronic version) of the Small Magellanic Cloud star cluster NGC~419, derived from ACS/WFC {\\it (bottom left)} and HRC {\\it (top right)} images in the F555W and F814W filters. The top left and bottom right panels zoom in onto the HRC image. Red-clump stars are marked with circles. At first glance, it is evident that they are quite homogeneous in their colours and luminosities, as expected for red-clump stars. However, when comparing their first Airy rings, one notices some quite subtle and {\\em systematic} differences in their brightnesses. Indeed, there are two kinds of red-clump stars: the most numerous and brighter, and a subsample (about 15\\% of the total) of fainter `secondary red-clump stars' (marked with red and green circles, respectively, in the electronic version). The luminosity difference between these groups is about 0.4~mag. There is also a $\\sim0.04$~mag difference in the mean colour which, however, is too small to be noticeable in the figure. The two kinds of red-clump stars are uniformly distributed across the HRC image. } \\label{fig_image} \\end{center} \\end{figure}  NGC~419 is a populous star cluster located to the east of the Small Magellanic Cloud (SMC)'s bar in a region relatively devoid of dust and free from contamination by the SMC field. There are two wonderful pairs of {\\sl HST} images of this cluster, taken in the F555W and F815W filters, originally obtained as part of programme GO-10396 (PI J. S. Gallagher) and now retrievable from the {\\sl HST} archive. They are shown in Figure~\\ref{fig_image}.  While the ACS/WFC images reveal the overall cluster structure, the ACS/HRC observations show the details of the cluster core.  \\begin{figure}[b] \\begin{center} \\includegraphics[width=0.6\\textwidth]{cmd_hrc.ps} \\caption{NGC~419 colour--magnitude diagram (CMD; {\\it left panel}), derived from the ACS/HRC data (Girardi et al. 2009). The error bars in the left panel are upper limits to the photometric errors. The two right-hand panels detail the red-clump and main-sequence turnoff, clearly showing their composite structure.} \\label{fig_cmd} \\end{center} \\end{figure}  The HRC images allow us to perform photometry of an impressive quality, and reveal at least two surprises: the presence of multiple main sequence turnoffs (MMSTOs), and a composite red clump which contains a pronounced faint extension (Figure~\\ref{fig_cmd}). Both features were noticed by Glatt et al. (2008). They tentativeley interpreted the faint red-clump structure as being caused by SMC field contamination. However, this explanation does not stand up to a simple star count in the neighbouring field selected from the ACS/WFC image (see Girardi et al. 2009): the expected number of field red-clump stars in the ACS/HRC area is just 4.5, while the faint red clump contains at least 50 objects (Figure~\\ref{fig_cmd}). Assuming a Poissonian distribution for the field stars, the probability ($P$) that these 50 stars are drawn from the SMC field alone is virtually zero ($P<10^{-9}$).  Detached binaries cannot give rise to this faint red clump either, since any combination of single stars will cause an extension of the red clump to brighter magnitudes.  What, then, is the dual red clump of NGC~419 made of? Girardi et al. (2009) show that it is caused by the presence of stars following two very different evolutionary paths. In the following, we will expand their arguments a little. There are, in fact, a few different ways of looking at dual red clumps, as detailed below. ",
        "conclusions": ""
    },
    "0910/0910.5145_arXiv.txt": {
        "abstract": "We have derived the quantum vacuum pressure $p_{\\rm vac}$ as a primary entity, removing a trivial and a gauge terms from the cosmological constant-like part (the zeroth term) of the effective action for a matter field. The quantum vacuum energy density $\\tilde{\\varrho}_{\\rm vac}$ appears a secondary entity, but both are of expected order. Moreover $p_{\\rm vac}$ and $\\tilde{\\varrho}_{\\rm vac}$ are dynamical, and therefore they can be used in the Einstein equations. In particular, they could dynamically support the holographic dark energy model as well as the ``thermodynamic'' one. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.2848_arXiv.txt": {
        "abstract": "We investigate the relationship between the linewidths of broad \\mgii\\ $\\lambda$2800 and \\hb\\ in active galactic nuclei (AGNs) to refine them as tools to estimate black hole (BH) masses.  We perform a detailed spectral analysis of a large sample of AGNs at intermediate redshifts selected from the Sloan Digital Sky Survey, along with a smaller sample of archival ultraviolet spectra for nearby sources monitored with reverberation mapping (RM).  Careful attention is devoted to accurate spectral decomposition, especially in the treatment of narrow-line blending and \\feii\\ contamination.  We show that, contrary to popular belief, the velocity width of \\mgii\\ tends to be smaller than that of \\hb, suggesting that the two species are not cospatial in the broad-line region.  Using these findings and recently updated BH mass measurements from RM, we present a new calibration of the empirical prescriptions for estimating virial BH masses for AGNs using the broad \\mgii\\ and \\hb\\ lines.  We show that the BH masses derived from our new formalisms show subtle but important differences compared to some of the mass estimators currently used in the literature. ",
        "introduction": "It is generally accepted that active galactic nuclei (AGNs) are powered by the release of gravitational energy from material accreted onto supermassive black holes (BHs).  The determination of BH mass (\\mbh) is crucial for understanding the AGN phenomena, the cosmological evolution of BHs, and even the coevolution of AGNs and their host galaxies.  Yet, for such distant objects, it is currently impossible to obtain direct measurement of \\mbh\\ using spatially resolved stellar or gas kinematics.  Fortunately, significant advances have been made in recent years from reverberation mapping (RM) studies of nearby Seyfert galaxies and quasi-stellar objects (QSOs; e.g., Wandel et al. 1999; Kaspi et al. 2000; Peterson et al. 2004).  First, the anti-correlation between the radius of the broad-line region (BLR) and the velocity width of broad emission lines for single objects supports the idea that the BLR gas is virialized and that its velocity field is dominated by the gravity of the BH (Peterson \\& Wandel 1999, 2000; Onken \\& Peterson 2002).  Second, the size of the BLR scales with the continuum luminosity (Kaspi et al. 2000, 2005), approximately as $R \\propto L^{0.5}$ (Bentz et al. 2006, 2009); the $R-L$ relation offers a highly efficient procedure for estimating the BLR size without carrying out time-consuming RM observations.  And third, the BH masses estimated by RM are roughly consistent (Gebhardt et al. 2000b; Ferrarese et al. 2001; Nelson et al. 2004; Onken et al. 2004) with the predictions from the tight correlation between \\mbh\\ and bulge stellar velocity dispersion established for inactive galaxies (the \\mbh--$\\sigma_\\star$ relation; Gebhardt et al.  2000a; Ferrarese \\& Merritt 2000).  These developments imply that we can estimate the BH mass in type 1 (broad-line, unobscured) AGNs by simple application of the virial theorem, \\mbh\\ = $fRv^2/G$, where $f$ is a geometric factor of order unity that depends on the geometry and kinematics of the line-emitting region, $R$ is the radius of the BLR derived from the AGN luminosity, and $v$ is some measure of the virial velocity of the gas measured from single-epoch spectra. The feasibility of obtaining $R$ and $v$ from single-epoch spectra enables \\mbh\\ to be estimated very efficiently for large samples of AGNs, especially for luminous quasars at higher redshift that typically exhibit only slow and small-amplitude variability (e.g., Kaspi et al. 2007), with the assumption that the virial relation is independent of redshift and can be extrapolated to higher luminosities and masses. In practice, for those AGNs that have measurements of $\\sigma_\\star$, $f$ is determined empirically by scaling the virial masses to the \\mbh--$\\sigma_\\star$ relation of inactive galaxies (e.g., Onken et al. 2004).  Implicit in this practice is the assumption---one open to debate (Greene \\& Ho 2006; Ho et al. 2008; Kim et al. 2008)---that active and inactive BHs should follow the same \\mbh--$\\sigma_\\star$ relation. The most widely used estimator for $v$ is the full width at half-maximum (FWHM) of the line.  Now, a large number of formalisms to estimate \\mbh\\ from single-epoch spectra have been proposed in the recent literature, using different broad emission lines optimized for different redshift regimes probed by (widely available) optical spectroscopy.  At low redshifts, the lines of choice are \\hb\\ (Kaspi et al. 2000; Collin et al. 2006; Vestergaard \\& Peterson 2006) or \\ha\\ (Greene \\& Ho 2005). At intermediate redshifts, \\mgii\\ $\\lambda$2800 is used (McLure \\& Jarvis 2002), while at high redshifts, one has to resort to \\civ\\ $\\lambda$1549 (Vestergaard 2002; Vestergaard \\& Peterson 2006). These formalisms are ultimately calibrated against RM masses based on the \\hb\\ BLR radius and linewidth measured from the variable (rms) spectra (Peterson et al. 2004).  Because the \\hb\\ linewidth is typically smaller in the rms spectra than in the single-epoch or mean spectra (Vestergaard 2002; Collin et al. 2006; Sulentic et al. 2006), some authors have proposed that the FWHM used in the \\hb-based formalisms should be further corrected to obtain unbiased \\mbh\\ estimates (Collin et al. 2006; Sulentic et al. 2006).  As for the \\mgii-based formalisms, because there are very few RM experiments of the \\mgii\\ line, they are either based on the RM data for \\hb\\ (e.g., McLure \\& Jarvis 2002; McLure \\& Dunlop 2004) or calibrated against the \\hb\\ formalisms themselves (e.g., Kollmeier et al. 2006; Salviander et al. 2007). A strong, underlying assumption is that \\mgii\\ and \\hb\\ are emitted from the same location in the BLR and have the same linewidth (see also Onken \\& Kollmeier 2008). In support of this assumption, some authors find that \\mgii\\ and \\hb\\ indeed have very similar linewidths (e.g., McLure \\& Jarvis 2002; McLure \\& Dunlop 2004; Shen et al. 2008; also cf. Salviander et al. 2007).  However, there are conflicting results in the literature: Corbett et al. (2003) claimed that \\mgii\\ is generally broader than \\hb, whereas Dietrich \\& Hamann (2004) came to an opposite conclusion.  Certainly, the most direct way to settle this issue is through direct RM of the \\mgii\\ line.  So far there are only two objects that have successful \\mgii\\ RM, NGC\\,5548 (Clavel et al. 1991; Dietrich \\& Kollatschny 1995) and NGC\\,4151 (Metzroth et al. 2006).  These studies tentatively suggest that \\mgii\\ responds more slowly to continuum variations than \\hb, implying that the \\mgii-emitting region is larger than that radiating \\hb.  Thus, there are still some important open questions regarding the robustness of \\mbh\\ measurements based on \\mgii.  What is the relation between the linewidths of \\hb\\ and \\mgii?   Are estimates of \\mbh\\ based on \\mgii\\ consistent with those based on \\hb?  These basic questions are critical for understanding the systematic uncertainties in studies of the cosmological evolution of BHs (cf. Shen et al. 2008; McGill et al. 2008; Denney et al. 2009a). To address the above questions, we perform a detailed comparison of the widths of the \\mgii\\ and \\hb\\ lines using single-epoch spectra for a large, homogeneous sample of Seyfert 1 nuclei and QSOs at intermediate redshifts culled from the Sloan Digital Sky Survey (SDSS; York et al. 2000). We further compare single-epoch \\mgii\\ linewidths with \\hb\\ linewidths measured from the rms spectra of AGNs with RM observations, finding systematic deviations between the two.  We present a recalibration of the \\mgii\\ virial mass estimator and compare our formalism with previous ones in the literature.  This paper adopts the following set of cosmological parameters: $H_{\\rm 0}$=70 km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_m $=0.3, and $\\Omega_{\\rm \\Lambda}$=0.7. ",
        "conclusions": "We investigate the relation between the velocity widths for the broad \\mgii\\ and \\hb\\ emission lines, derived from FWHM measurements of single-epoch spectra from a homogeneous sample of 495 SDSS Seyfert 1s and QSOs at $0.45 < z <0.75$. Careful attention is devoted to accurate spectral decomposition, especially in the treatment of narrow-line blending and \\feii\\ contamination. We find that \\mgii\\ FWHM is systematic smaller than \\hb\\ FWHM, such that FWHM(\\mgii)\\,$\\propto$\\,FWHM(\\hb)$^{0.81\\pm 0.02}$.  Using 29 AGNs that have optical RM data and usable archival UV spectra, we then investigate the relation between single-epoch \\mgii\\ FWHM and rms \\hb\\ \\sigline\\ (line dispersion), a quantity regarded as a good tracer of the virial velocity of the BLR clouds emitting the variable \\hb\\ component. We find that, similar to the situation for the FWHM of single-epoch \\hb, single-epoch \\mgii\\ FWHM is unlikely to be linearly proportional to rms \\hb\\ \\sigline.  The above two findings suggest that a major assumption of previous \\mgii-based virial BH mass formalisms---that the \\mgii-emitting region is identical to that of $\\hb$---is problematic.  This finding and the recent updates of the reverberation-mapped BH masses (Peterson et al. 2004; Denney et al. 2006, 2009b; Metzroth et al. 2006; Bentz et al. 2007; Grier et al. 2008) motivated us to recalibrate the \\mbh\\ estimator based on single-epoch \\mgii\\ spectra.  Starting with the empirically well-motivated BLR radius--luminosity relation and the virial theorem, $\\mbh \\propto L^\\beta {\\rm FWHM}^\\gamma$, we fit the reverberation-mapped objects in a variety of different ways to constrain $\\beta$ and $\\gamma$.  For all the strategies we have considered, $\\beta$ has a well-defined value of $\\sim$0.5, in excellent agreement with the latest BLR radius--luminosity relation (Bentz et al. 2006, 2009), whereas $\\gamma \\approx 1.5 \\pm 0.5$, which is marginally in conflict with the canonical value of $\\gamma = 2$ normally assumed in past studies. Performing a similar exercise for \\hb\\ yields $\\mbh \\propto L^{0.5} {\\rm FWHM(\\hb)}^{1.09\\pm0.22}$, which again significantly departs from the functional forms used in the literature. The 1 $\\sigma$ uncertainty (scatter) is of the order of 0.3 dex relative to the RM-based masses for the \\hb\\ estimator, and $\\sim$0.4 dex for the \\mgii\\ estimator. Using the same data set, the scatter of our \\hb\\ mass scaling relation is reduced by  0.1 dex over that of Vestergaard \\& Peterson (2006), indicating improvement in the internal scatter.  We use the SDSS database to compare our new \\mbh\\ estimators with various existing formalisms based on single-epoch \\hb\\ and \\mgii\\ spectra.  BH masses derived from our \\mgii-based mass estimator show subtle but important deviations from many of the commonly used \\mbh\\ estimators in the literature.  Most of the differences stem from the recent recalibration of the masses derived from RM.  Researchers should exercise caution in selecting the most up-to-date \\mbh\\ estimators, which are presented here."
    },
    "0910/0910.2765_arXiv.txt": {
        "abstract": "We find statistically significant spatial correlations between the arrival directions of the highest energy cosmic rays (HECRs) observed by the Akeno Giant Air Shower Array (AGASA) and large-scale structure of galaxies observed by Sloan Digital Sky Survey (SDSS) in the redshift ranges of $0.006 \\leq z < 0.012$ and $0.012 \\leq z < 0.018$ at angular scale within $\\sim 5^{\\circ}$. This result supports a hypothesis that the sources of HECRs are related to galaxy distribution even in the northern sky, which has been already indicated by Pierre Auger Observatory in the southern sky. We also investigate the dependency of the correlation on the absolute magnitude, color, and morphology of the galaxies. For galaxies with $0.006 \\leq z < 0.012$, the correlation tends to be stronger for luminous and red galaxies. Based on these results, we discuss plausible HECR sources and constraint on Galactic magnetic field. ",
        "introduction": "\\label{introduction}  Although the sources of the highest energy cosmic rays (HECRs) have been poorly known, recent observations begin to unveil a part of the nature of HECR origin. The first indication to HECR sources was reported by the Akeno Giant Air Shower Array (AGASA). The AGASA reported the anisotropic distribution of cosmic rays (CRs) above $4 \\times 10^{19}$ eV at small angular scale comparable with their angular uncertainty to determine arrival directions of primary CRs \\cite{takeda99}. The anisotropic feature has been interpreted as the existence of point-like extragalactic sources. An important step to their origin is the positional correlation between the arrival directions of HECRs with energies above $\\sim 6 \\times 10^{19}$ eV and nearby astrophysical objects reported by the Pierre Auger Observatory (PAO) \\cite{abraham07,abraham08}. These results strongly suggested that HECRs should be of extragalactic origin. Several groups have also analyzed the PAO data by using catalogs of different astrophysical objects and then have confirmed the correlation with matter (or galaxy) distribution of local Universe \\cite{kashti08,george08,ghisellini08,takami09c}. Although the same analysis of new PAO data as Refs. \\cite{abraham07,abraham08} shows the decreases of the significance level of the correlation \\cite{hague09}, the correlation with the large-scale structure is still positive \\cite{aublin09}.   The PAO observes HECRs from the southern sky, since it is located in Argentina. If HECR sources are extragalactic objects as suggested by the PAO, the correlation of HECRs and galaxy distribution is also expected in the northern sky. However, such correlation has never been clearly reported, though the correlation of HECRs with the supergalactic plane was reported in the northern sky \\cite{stanev95,uchihori00}. We tested the correlation of the AGASA data with galaxies of Infrared Astronomical Satellite Point Source Redshift Survey (IRAS PSCz) \\cite{saunders00}, but a significantly positive correlation could not be found \\cite{takami09c}. There are two reasons why we have not found such correlation when we assume that the correlation exists in the northern sky. One is the small exposure of the AGASA relative to the PAO. In the period of Refs. \\cite{abraham07,abraham08}, the total exposure of the PAO is $\\sim$ 7 times as large as that of the AGASA. Thus, the AGASA may not have the number of highest energy (HE) events large enough to find correlation. The other is the completeness of a galaxy catalog used in \\cite{takami09c}. In other words, it is the possibility that the IRAS PSCz catalog of galaxies does not reflect the large-scale structure of galaxy distribution in local Universe.   In this study, we focus the latter reason, because the former problem cannot be solved by us without accumulating more number of HECR events. We investigate the correlation between HECR events observed in the northern sky and galaxy distribution which has higher completeness than the IRAS PSCz catalog. For this purpose, we adopt galaxies observed by Sloan Digital Sky Survey (SDSS) \\cite{york00}. The SDSS is a survey project of galaxies to investigate the large-scale structure of the Universe in optical bands and therefore can observe faint galaxies down to $m_r \\sim 20.0$ where $m_r$ is the $r$-band apparent magnitude of galaxies. Thus, SDSS galaxies better trace the structure of the local Universe. As HECR events observed in the northern sky, we adopt the AGASA data published in Ref. \\cite{hayashida00}.   This paper is laid out as follows: in Section \\ref{methods}, we explain the data samples of HECR events and galaxies in detail. Also, a statistical method which we adopt to test the correlation between them is described. In Section \\ref{results}, we calculate cross-correlation functions defined in Section \\ref{methods} between the AGASA events and SDSS galaxies with several redshift ranges. We also investigate the dependencies of the correlation on the $r$-band absolute magnitude, color, and morphology of galaxies. These give us useful information to search for the HECR sources because the features of the host galaxies reflect the nature of objects contained in the hosts. Section \\ref{discussion} is devoted to discuss and interpret the results obtained in Section \\ref{results}. Finally, we conclude this study in Section \\ref{conclusion}. ",
        "conclusions": "\\label{conclusion}  In this study, we investigated the correlation between the arrival directions of HECRs detected in the northern sky (the AGASA data) and galaxy distribution observed by the SDSS. The SDSS can observe faint galaxies, and therefore SDSS galaxies are a good sample to investigate the correlation with the large-scale structure of matter in local Universe. We found the positive correlation of the AGASA events with SDSS galaxies within the angular scale of $\\sim 5^{\\circ}$ with the redshift ranges of $0.006 \\leq z < 0.012$ and $0.012 \\leq z < 0.018$ for the first time. This result supported a hypothesis that the sources of HECRs are related to  galaxy distribution even in the northern sky, which has been already indicated by several analyses of the PAO data in the southern sky \\cite{abraham07,abraham08,kashti08,george08,ghisellini08,takami09c}. We also tested the dependencies of the correlation signals on the r-band absolute magnitude, color, and morphology of the galaxies, since these information is generally useful to constrain HECR sources involved in the host galaxies. For galaxies with $0.006 \\leq z < 0.012$, we found that the AGASA events correlated with luminous ($M_r \\leq -19$) and red ($M_g - M_r > 0.6$) galaxies, and the luminous galaxies around the AGASA events were relatively red and morphologically late-type. Such dependencies were not clear for galaxies with $0.012 \\leq z < 0.018$. This was possible because the flux of HECRs from a sources with this redshift range is smaller than for $0.006 \\leq z < 0.012$ and/or the trajectories of HECRs are more deflected than for sources with $0.006 \\leq z < 0.012$.   We discussed possible source candidates based on our results for galaxies with $0.006 \\leq z < 0.012$ in Section \\ref{discussion}. Unfortunately, the interpretation of the results is ambiguous at present because of uncertainty of relations between HECR source candidates and the nature of their host galaxies, and our poor knowledge on intervening magnetic fields. Under the limitations, possible interpretation is summarized as follows: Morphology dependence of galaxies correlating with the AGASA events is positive for GRBs and magnetars, but the colors of these galaxies are relatively red, contrary to relatively blue colors expected for these objects. Note that the discussion on the colors was only qualitative. More detail and quantitative discussion is required to confirm/reject GRB or magnetar origin of HECRs. On the other hand, HECR generation by non-thermal phenomena related to AGNs might be difficult from results in this study because the morphologies of luminous galaxies near the AGASA events are relatively late-type, though such phonomena are strongly related to galaxies with early-type galaxies.   The interpretation above is based on the assumption that the trajectories of HECRs are not largely deflected by IGMF. When IGMF is strong, fake correlation could occur \\cite{kotera08,ryu09}. The correlation between the arrival directions of HECRs and the positions of their sources is reduced, though the correlation with matter distribution is conserved. In this case, it is not possible to confirm HECR sources by the nature of galaxies correlating with HECR events.   The angular scale of the positive correlation enabled us to constrain intervening magnetic fields under the assumption that we see the correlation between HE protons and their sources. The angular scale estimated in this study, $< 5^{\\circ}$, was compared with theoretical predictions by Ref. \\cite{takami09f}. The angular scale was consistent with predictions based on bisymmetric spiral field models of GMF. Since the larger angular deflections of HE protons are predicted for axisymmetric GMF models, it might be difficult to realize the axisymmertic models.  In order to investigate correlation signals between HECRs and matter distribution/their sources more accurately, we need more number of HECR events. It was not clear whether the correlation features found in this study for galaxies with $0.006 \\leq z < 0.012$ are common in the Universe. This can be checked by increasing the observed number of HECRs. We can also understand intervening magnetic fields and the nature of HECR sources. Ref. \\cite{takami08a} predicted the spatial correlation between HE protons and their sources in local Universe even taking a realistically structured IGMF model into account, which was clearly visible by accumulating $\\sim 200$ events above $\\sim 6 \\times 10^{19}$ eV for the number density of HECR sources of $10^{-5}$ Mpc$^{-3}$. Data accumulation in the northern sky is in progress. The total exposure of the Telescope Array have already reached 75\\% of the AGASA exposure \\cite{taketa09} and is expected to reach that of the AGASA in this year. Thus, their data is quite useful to test the correlation and the nature of the host galaxies of HECR sources. Also, two projected HECR observatories with extremely large effective area, Extreme Universe Space Telescope (JEM-EUSO) \\cite{ebisuzaki09} and the northern site of the PAO \\cite{harton09}, enable us not only to clarify the correlation but also to unveil the positions of HECR sources in the northern sky of local Universe.  \\ack We are grateful to K.Maeda and I.Kayo for useful discussion. This work is supported by Grants-in-Aid for Scientific Research from the Ministry of Education, Culture, Sports, Science and Technology (MEXT) of Japan through No.21840019 (H.T.) and No.19104006 (K.S.). The work of T.N. is supported by the JSPS fellow. This work is also supported by World Premier International Research Center Initiative (WPI Initiative), from the MEXT of Japan.   Funding for the SDSS and SDSS-II has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Science Foundation, the U.S. Department of Energy, the National Aeronautics and Space Administration, the Japanese Monbukagakusho, the Max Planck Society, and the Higher Education Funding Council for England. The SDSS Web Site is http://www.sdss.org/.  The SDSS is managed by the Astrophysical Research Consortium for the Participating Institutions. The Participating Institutions are the American Museum of Natural History, Astrophysical Institute Potsdam, University of Basel, University of Cambridge, Case Western Reserve University, University of Chicago, Drexel University, Fermilab, the Institute for Advanced Study, the Japan Participation Group, Johns Hopkins University, the Joint Institute for Nuclear Astrophysics, the Kavli Institute for Particle Astrophysics and Cosmology, the Korean Scientist Group, the Chinese Academy of Sciences (LAMOST), Los Alamos National Laboratory, the Max-Planck-Institute for Astronomy (MPIA), the Max-Planck-Institute for Astrophysics (MPA), New Mexico State University, Ohio State University, University of Pittsburgh, University of Portsmouth, Princeton University, the United States Naval Observatory, and the University of Washington.    \\clearpage"
    },
    "0910/0910.5659_arXiv.txt": {
        "abstract": "Hamiltonian perturbation theory is used to analyse the stability of $f(R)$ models. The Hamiltonian equations for the metric and its momentum conjugate are written for $f(R)$ Lagrangian in the presence of perfect fluid matter. The perturbations examined are perpendicular to $R$. As perturbations are added to the metric and momentum conjugate to the induced metric instabilities are found, depending on the form of $f(R)$. Thus the examination of these instabilities is a way to rule out certain $f(R)$ models. ",
        "introduction": "The question of dark energy has been at the heart of cosmology since the discovery of accelerating expansion of the universe \\cite{riess98}. The traditional picture of general relativity with ordinary relativistic or non-relativistic matter in homogeneous and isotropic universe meets severe problems when accommodating it to current cosmological observations. The conflicting observational evidence comes mainly from supernova light curves \\cite{riess98,perlmutter99}, CMB anisotropies \\cite{Bennett2003,Netterfield2001} and large scale structures \\cite{Perlmutter1999,Bunn1996}. This has lead to several suggested remedies. Perhaps the most popular way is to add some non-conventional matter to the universe. Among these the simplest possibility is no doubt to use the cosmological constant. A review of the subject can be found in \\cite{peebles2002}. In any case, the key aspect is the negative pressure of the new matter which boosts the expansion of the universe. Other considerations include more general distribution of matter, {\\it i.e.} non-homogeneous or non-isotropic universe (see {\\it e.g.} \\cite{Alnes2005}).  Besides these two, a lot of effort has been put into studies on generalizations and modifications of General Relativity. For example metric-affine theories (see {\\it e.g.} \\cite{Sotiriou2006}), scalar-tensor theory (see {\\it e.g.} \\cite{faraoni04,Bransdicke1961}), brane-world gravity (see {\\it e.g.} \\cite{Maartens2003}) and more general Lagrangians have been considered. In the present paper we are especially interested in $f(R)$ gravity models in which the Einstein-Hilbert action is replaced by a function of the curvature scalar $R$ \\cite{Carroll2003,Carroll2004,Meng2003, Allemandi2004,Sotiriou2008}. None of these modifications is free of problems and this is indeed the case of $f(R)$ gravity as well. As for any model, the cosmological observations issue some constraints (see {\\it e.g.} \\cite{Tsujikawa2007,Starobinsky2007}) as do the observations in the solar system (see {\\it e.g.} \\cite{Chiba2003, Hu2007,Magnano93,multamaki2008,Henttunen2007}). The opinions are still divided on the viability of $f(R)$ theories of gravity. There are numerous approving studies (see {\\it e.g.} \\cite{Faraoni2006,Olmo2006}) as well as sceptical ones (see {\\it e.g.} \\cite{Faulkner2006,Erickcek2006}).   As the actual universe is not homogeneous and isotropic but contains local perturbations, additional challenges for $f(R)$ theories emerge from stability analysis \\cite{Dolgov2003,Soussa2003,Faraoni2005}. An acceptable cosmological model has to be stable against perturbations in the metric and the mass distribution. However, stability analysis is customarily done only in the direction of $R$, {\\it i.e.} only curvature perturbations are considered. This is motivated in particular in the case of General Relativity, where the relation between space-time curvature and the matter density is a simple one: the trace of Einstein equations imply $R\\propto \\rho$. This in turn implicates a simple and direct relation between the perturbations in matter and curvature. This is not the only possibility. In a $f(R)$ model the model the relation is more complicated due to appearance of function $f(R)$ and higher derivative terms in the field equations. The phase space is considerably larger and metrics corresponding a given matter distribution ambiguous. The physical acceptability, however, of a model requires general stability; also stability against perturbations which keep curvature constant, perpendicular to $R$.  The Hamiltonian formulation of general relativity has been around since the work of Arnowitt, Deser and Misner \\cite{Arnowitt1962}\\footnote{The ideas were first seen in the long out of print {\\it Gravitation: an introduction to current research}. The authors have later on released the article on ArXiv as cited.}. Hamiltonian formulation has also surfaced in the works of Ashtekar \\cite{ashtekar87}. The first papers on the subject often neglected the boundary terms, however, later works have clarified these details ({\\it e.g.} \\cite{Brown1992, Hawking1995}). Hamiltonian formulation has not received too much interest in contemporary papers. In particular and to our knowledge the use of Hamiltonian formulation on perturbations of $f(R)$ theories has not been studied so far. The main interest has been in specific choices for the function $f(R)$.  In the present paper we look into perturbations using Hamiltonian formalism of $f(R)$ theories. While the technique has not yet been applied to general $f(R)$ theories with perturbations it is a useful tool in studying the stability of $f(R)$ models: with it is simple to study perturbations perpendicular to $R$. As in classical mechanics the Hamiltonian is written as a functional of the fields and their canonical momenta. However, in a geometric theory like General Relativity and $f(R)$ theories, some complications appear due to constraints between field components. The two main aspects of the canonical Hamiltonian formalism are that the field equations are of the first order in the time derivatives and that time is distinguished from other coordinates. For writing the Hamiltonian equations we must thus foliate the region of space-time with space-like hypersurfaces. Finally the resulting field equations for the perturbations are then analyzed for instabilities. The conventions and details of the formalism can be found in \\cite{poisson2004}.  The paper is organized as follows. In section \\ref{hamform} we write the Hamiltonian field equations. The 3+1 decomposition and foliation of the space-time are also presented. The first order perturbations are added to the system in section \\ref{1order}. We also take a look at second order perturbations in section \\ref{2order}. In section \\ref{concl} we summarize our results. ",
        "conclusions": "\\label{concl}  Traditionally the stability analysis is performed in the Lagrangian formalism and the analysis parallel to $R$ has been carried out before in several papers ({\\it e.g.} \\cite{Dolgov2003}). Many of the interesting $f(R)$ models have been found to be inherently unstable in the past \\cite{Kainulainen20071,Faraoni2005}. However, stability analysis has not yet been used to the full extent as long as the studies concentrate on curvature perturbations only. By using the Hamiltonian instead of Lagrangian formulation we examined the large scale cosmological perturbations perpendicular to $R$ with non-relativistic matter. These perturbations are fairly easy to examine with the Hamiltonian formulation. The found instabilities are noticeably different to those of previous works ({\\it e.g.} \\cite{Faraoni2005}). Because of the constraint $\\dot R=0$ diagonal perturbations of metric and conjugate momentum are related to each other. The temporal part of the metric showed up to be constrained by the conjugate momentum one. Moreover, the spatial part of the metric is forced to vanish. If these constraints were not satisfied we would have non-trivial perturbations of non-diagonal elements of the metrics like $g_{0i}$.  The perturbations of the momentum conjugate turn out to be the most interesting ones. The equation depends explicitly on the form of the function $f(R)$. Some choices of $f(R)$ lead clearly unstable cosmological model, but seemingly not all. We have studied some well-known $f(R)$ functions and find them unstable. Albeit the physical interpretation of the perturbation momentum conjugate is unfortunately not as clear as that of the metric perturbations, equation \\eqref{pab} demonstrates the relation between the momentum conjugate and the extrinsic curvature. In the 3+1 decomposition the intrinsic curvature $\\tilde R$ defines how the hypersurface is curved whereas the extrinsic curvature defines how each slice is curved relative to the enveloping space-time.  As the perturbations were not well-behaved in this context further studies would be relevant in order to find the limits of these constraints. Fruitful directions would likely to be investigating the effects of more other types of matter. Also, it would be prudent to examine the case where the metric can include shift ({\\it i.e.} $N_a\\neq 0$). It is clear from the form of \\eqref{hamden} that such a generalization would affect the following equations of motion deeply as the last term would be non-zero. This is understandable as the metric would now include spatio-temporal elements. It is also possible to study more general theories with the Lagrangian depending also on for example $R_{\\mu\\nu}R^{\\mu\\nu}$ or Gauss-Bonnet term.  It appears that with Hamiltonian formulation of perturbations can be used to constrain the spectrum of cosmologically acceptable $f(R)$ theories. While there are several physical arguments to judge the $f(R)$ theories like cosmological observations and solar system behaviour, stability analysis is one important tool to rule out ill-behaved models out of numerous possible modified theories of gravity. With continued studies it is possible to find the ones best describing the observed behaviour of the universe.  \\subsection*"
    },
    "0910/0910.0774_arXiv.txt": {
        "abstract": "Compared to starburst galaxies, normal star forming galaxies have been shown to display a much larger dispersion of the dust attenuation at fixed reddening through studies of the IRX-$\\beta$ diagram (the IR/UV ratio ``IRX'' versus the UV color ``$\\beta$''). To investigate the causes of this larger dispersion and attempt to isolate second parameters, we have used GALEX UV, ground-based optical, and Spitzer infrared imaging of 8 nearby galaxies, and examined the properties of individual UV and 24~$\\mu$m selected star forming regions. We concentrated on star-forming regions, in order to isolate simpler star formation histories than those that characterize whole galaxies. We find that 1) the dispersion is not correlated with the mean age of the stellar populations, 2) a range of dust geometries and dust extinction curves are the most likely causes for the observed dispersion in the IRX-beta diagram 3) together with some potential dilution of the most recent star-forming population by older unrelated bursts, at least in the case of star-forming regions within galaxies, 4) we also recover some general characteristics of the regions, including a tight positive correlation between the amount of dust attenuation and the metal content. Although generalizing our results to whole galaxies may not be immediate, the possibility of a range of dust extinction laws and geometries should be accounted for in the latter systems as well. ",
        "introduction": "\\label{sec:introduction} Star formation is a fundamental process that governs the evolution of the baryonic matter in galaxies. It can deeply affect the host galaxy, its properties and appearance. The newly formed populations change the galaxy's spectral energy distribution (SED); the metals formed in the most massive stars pollute the interstellar medium and, in the presence of mechanical feedback from episodes of intense star formation, can be ejected from the galaxy and pollute the intergalactic medium. It is important to accurately measure star formation both as a function of spatial position within a galaxy and as a function of a galaxy's temporal evolution. The cosmic star formation rate density as a function of redshift \\citep{madau1996a,hopkins2006b} has become a classic comparison benchmark, and a challenge, for both semi-analytic and hydrodynamical models of structure formation and evolution \\citep[e.g.][]{katz1996a,hopkins2006a,stinson2006a,cattaneo2007a}. The youngest stellar populations, whose bolometric energy output is dominated by short-lived massive stars, emit the bulk of their energy in the restframe ultraviolet (UV). Since the pioneer work of \\cite{donas1984a} the UV has become one of the main star formation rate (SFR) tracers used in modern Astronomy \\citep{kennicutt1998a}. This regime, which has been previously difficult to observe in the local Universe, becomes the wavelength region ``par excellence'' to investigate star formation in galaxies at intermediate and high redshift. At these distances, the restframe UV is shifted to optical and near-infrared wavelengths, where current detectors afford the highest performance in terms of both sensitivity and angular resolution. Thus, it is no surprise that so far galaxy investigations in the redshift range z$\\sim$2-7 have been dominated by restframe UV observations \\citep{giavalisco2004a,bouwens2005a,mobasher2005a,sawicki2006a}. UV observations are the main canvass from which our current understanding of cosmic star formation is built.  For over a decade now, studies of galaxies at cosmological distances have employed the ``starburst attenuation curve'' \\citep{calzetti1994a,calzetti2000a,meurer1999a} -- see next paragraph for a description of the physical conditions leading to such a law and equation 4 from \\cite{calzetti2000a} -- to correct their restframe UV measurements for the effects of dust extinction \\citep[e.g.][to name a few]{steidel1996a,steidel1999a,giavalisco2004a,daddi2005a,daddi2007a}. This curve is a powerful tool, in that it only requires information on the UV colors (or UV spectral slope) of a galaxy to provide an extinction correction, and recover the intrinsic UV flux. Assuming this curve, the typical UV extinction correction at redshift 3 is a factor $\\sim5$ in luminosity \\citep{steidel1999a}. However, in recent years, evidence has been accumulating that the starburst attenuation curve is not applicable to all galaxies \\citep{buat2002a,buat2005a,noll2009a} and that the UV extinction is dependent on the age of the stellar populations \\citep{cortese2008a}. More quiescently star-forming galaxies show characteristics in their UV SED that do not lend themselves to the same extinction ``treatment'' as more active starbursts. Applying the starburst attenuation curve to these more quiescent galaxies leads to up to an order of magnitude overestimate of the UV luminosity (and of the SFR). This result clearly represents a problem as deeper observations probe past the most powerful starburst to more normal galaxies with relatively lower SFR at increasingly higher redshifts. Along these lines, current determinations of UV luminosity functions at cosmological distances may be viewed with some suspicion, as all of them need some form of extinction correction \\citep{steidel1999a,giavalisco2004a,sawicki2006a,teplitz2006a}.  One of the prescriptions of the starburst attenuation curve is based on a very simple, fully-empirical result. The ratio L(IR)/L(UV) is strongly correlated with the measured UV spectral slope $\\beta$ (or the UV color) in local, UV-bright starburst galaxies \\citep{meurer1999a,calzetti2000a}. The L(IR)/L(UV) ratio is a measure of the total dust opacity that affects the UV emission in a galaxy, or more generally, in a stellar system; the UV light absorbed by dust is re-emitted in the IR, independently of the details of the star-dust distribution, the dust properties, the metallicity, the stellar initial mass function, or the star formation history of the system \\citep{gordon2000a,buat2005a}. Meurer et al.'s result, commonly called the IRX-$\\beta$ relation, together with a host of previous results that showed strong correlations between $\\beta$ and a number of dust reddening measures, form the basis for an extinction correction prescription that obviously applies to local (and possibly distant) starburst galaxies \\citep{calzetti1994a,calzetti1996a,calzetti1997a,calzetti2001a}. The properties of the starburst attenuation curve imply specific conditions for the distribution of dust and stars within the starburst region \\citep{calzetti1994a,calzetti1996a,witt2000a,charlot2000a,calzetti2001a}. The geometry that can account for all observed properties at UV, optical, near-IR, and far-IR wavelengths is that of a dust distribution foreground to the stellar population responsible for the bulk of the UV emission. The term foreground is here loosely used to indicate any geometry that puts the stars in the center of a shell-like dust distribution, and does not make any inference on the dust distribution (which can be homogeneous or clumpy). Starbursts provide a mechanism for creating such specific dust-star geometry: massive star wind and supernova feedback can create expanding bubbles of hot gas that sweep away a significant fraction of the dust and gas that resides in the starburst site itself. Within this scenario, the gas and dust end up at the ``edges'' of the starburst site thus creating an inhomogeneously distributed ``shell''. In addition, the starbust extinction curve also requires a specific type of dust grains lacking the 2175\\AA\\ bump \\citep{gordon1997a} like that found in the SMC bar \\citep{gordon1998a,gordon2003a}.  The starburst attenuation curve may thus not apply to a generic galaxy, as the coherent feedback mechanism present in starbursts may be missing in other galaxies. A number of studies have confirmed that more quiescent star-forming galaxies, although they generally trace a correlation in the IRX-$\\beta$ plane, are also up to an order of magnitude less obscured than starbursts for similar $\\beta$ values, and can have a larger dispersion, with a factor up to $\\sim$5 larger spread around the mean correlation \\citep{buat2002a,buat2005a,bell2002a,kong2004a,gordon2004a,seibert2005a,calzetti2005a,boissier2007a,dale2007a,munoz2009a}. UltraLuminous Infrared Galaxies (ULIRGs) also deviate from the starburst IRX-$\\beta$ correlation \\citep{goldader2002a}, and are located above it, i.e., opposite to the star-forming galaxies. ULIRGs are dustier, more compact, and denser systems than the UV-bright starbursts; therefore, their location in the IRX-$\\beta$ diagram is plausibly accounted for by an ``homogeneous'' mixture of dust and stars \\citep{calzetti2001a}. ULIRGs likely represent less of a concern for the use of the UV emission as a SFR tracer, as in general they will not be detectable at restframe UV wavelengths in cosmological samples. Quiescent star-forming galaxies are a far larger concern, as they are relatively bright in the UV and become increasingly detected as surveys go deeper and, by the nature of the luminosity function, are more numerous than starbursts. Alas, for these galaxies, the next step in the investigation, understanding and quantifying the nature of the deviation from the starbursts' IRX-$\\beta$ correlation, has so far eluded any attempt. \\cite{bell2002a} suggested that a combination of variations in the dust geometry and the characteristics (age, etc.) of the stellar populations could be the reason for the deviation observed in the HII regions of the Large Magellanic Cloud. The sparse data available, however, prevented further insights into the problem. Similar limitations were encountered by \\cite{gordon2004a} in their analysis of the star-forming regions in M~81. \\cite{kong2004a}, using a non-homogeneous UV dataset, suggested that the deviations of a large sample of quiescent star-forming galaxies from the starburst relation could be parametrized by the b-parameter, the ratio of the current to past average star formation rate. However, neither \\cite{seibert2005a}, using the more homogeneous GALEX All-sky Imaging Survey sample plus IR data from IRAS, nor \\cite{johnson2007b}, using SDSS data, Spitzer, and spectroscopic observations, confirmed Kong et al.'s suggestion. Some of these works have also produced ``mean'' obscuration curves for quiescent star-forming galaxies, but the intrinsic spread of the data around these mean curves is more than an order of magnitude, versus $\\sim$2-3 typical of the starbursts. This suggests the existence of a ``second parameter'' for the star-forming galaxies \\citep{kong2004a}, which needs to be identified and understood if the UV is going to be used as a tracer of SFR.  In this paper, we investigate the possible reason for the deviation from the starburst IRX-$\\beta$ relation concentrating on sub-kpc star forming regions within galaxies, which provide simpler samples of stellar population than whole galaxies. We use GALEX observations to measure the UV luminosity, Spitzer ones to derive the total infrared luminosity and optical ones to estimate the age.  In section 2 we present the sample of galaxies selected for this study, in section 3 we present the UV, IR and optical observations, in section 4 we present the data processing and photometry methods carried out, in section 5 we present how we model the star forming regions, in section 6 we proceed on to test the validity of the age as the ``second parameter'', in section 7 we test different extinction laws, in section 8 we test the metallicity, we discuss the results presented in section 9 and we finally conclude in section 10. ",
        "conclusions": "We have performed an analysis of {\\bf over 300} star forming regions in 8 local quiescent star forming galaxies to garner new insights on the IRX-$\\beta$ relation by using Spitzer infrared, GALEX ultraviolet and ground based optical broad band and narrow-band H$\\alpha$ data. Specifically we investigate the deviations of star forming regions in normal star forming galaxies from the locus identified by the starburst attenuation curve, in hope of obtaining information on similar deviations by normal star forming galaxies.  The mean age of the stellar population, as traced by the $U-B$ color and by the equivalent width of H$\\alpha$, is shown to be relatively uncorrelated or only weakly correlated with the perpendicular distance of the data from the starburst attenuation curve in the IRX-beta diagram. Stellar population ages have been suggested as a potential culprit for the deviations in the case of normal star-forming galaxies \\citep[e.g.][via the b-parameter]{kong2004a}, but we do not find evidence for such trend in HII knots of galaxies. Our results are in agreement with similar results, obtained for whole galaxies, by \\cite{seibert2005a} and \\cite{johnson2007b}, but disagree with recent findings by Dale et al. (2009), also derived for whole galaxies.  We find that a range of dust extinction/attenuation curves and dust/star geometries are actually required to account for the location on the IRX-beta plot of most of our HII knots. Specifically, we find that the starburst attenuation curve and the SMC extinction with foreground geometry bracket the location of most of the data points, not only on the IRX-beta diagram but on other plots reporting colors and other properties of the regions. For the reddest (in UV color) regions, we need to include contributions to the observed colors and fluxes of underlying older ($\\sim$ 300 Myr) stellar populations to account for the observed UV and $U-B$ colors. The mass ratio between the young (ionizing) stellar population and the older one is about 1/5.  Individual star forming regions and integrated galaxies do not populate the same locus in the IRX-$\\beta$ diagram. Indeed, the former have a bluer UV color and are more tightly correlated. The most likely explanation is that older stellar populations up to $1-2\\times10^9$ years old, which have a redder UV color redden the observed UV color of star forming galaxies. To accurately determine the star formation rate, it is crucial to take into account this older population.  Finally, a strong correlation is found with the metallicity, in the sense that the more metal rich regions tend to display redder UV colors and larger IR/UV values, as expected metallicity and dust content correlate with each other."
    },
    "0910/0910.5221_arXiv.txt": {
        "abstract": "We calculate how the relic density of dark matter particles is altered when their annihilation is enhanced by the Sommerfeld mechanism due to a Yukawa interaction between the annihilating particles. Maintaining a dark matter abundance consistent with current observational bounds requires the normalization of the s-wave annihilation cross section to be decreased compared to a model without enhancement. The level of suppression depends on the specific parameters of the particle model, with the kinetic decoupling temperature having the most effect. We find that the cross section can be reduced by as much as an order of magnitude for extreme cases. We also compute the $\\mu-$type distortion of the CMB energy spectrum caused by energy injection from such Sommerfeld-enhanced annihilation. Our results indicate that in the vicinity of resonances, associated with bound states, distortions can be large enough to be excluded by the upper limit $\\vert\\mu\\vert\\leq9.0\\times10^{-5}$ found by the COBE/FIRAS experiment. ",
        "introduction": "If dark matter annihilates, the byproducts of the annihilation (positrons, neutrinos, gamma-rays, etc.) can leave non-gravitational signatures that, if observed, would be crucial for clarifying the nature of dark matter.  In recent years, a number of observations have highlighted anomalies that might be explained by invoking dark matter annihilation in our Galactic halo. Among them are: (i) an anomalous abundance of positrons in cosmic rays above $10$~GeV according to the PAMELA experiment \\cite{Adriani-09} confirming and extending previous measurements by experiments such as AMS-01 \\cite{-07}; (ii) an excess of microwave emission from the galactic center as measured by the WMAP experiment, known as the ``WMAP haze'' \\cite{Hooper-Finkbeiner-Dobler-07} (iii) the apparent excess of diffuse galactic gamma-rays with energies above 1GeV inferred from observations by the EGRET satellite \\cite{deBoer-05}; (iv) analyses on the balloon experiments ATIC and PPB-BETS \\cite{Chang-08,Torii-08} have reported an excess in the total flux of electrons and positrons in cosmic rays. For the last two anomalies, we note that recent observations by FERMI seem inconsistent with the claimed gamma-ray excess in the EGRET data \\cite{Porter--09}, and that the excess in the electron and positron flux is smaller than previously thought \\cite{Abdo-09}. The latter is actually consistent with a modified cosmic ray propagation model that does not require additional primary sources of electrons and positrons \\cite{Grasso-09}.  Although other astrophysical sources could explain these anomalies (see for example \\cite{Hooper-Blasi-DarioSerpico-09,Yuksel-Kistler-Stanev-09,Profumo-08} for an explanation for the PAMELA excess based on particle acceleration by pulsars, and \\cite{Fujita-09,Shaviv-Nakar-Piran-09} for one based on supernova remnants), dark matter annihilation offers an attractive solution. Large annihilation rates are needed, however, to explain the observations. In particular for the PAMELA data, the required annihilation rate is typically a few orders of magnitude larger than the value obtained by assuming the standard cross section inferred from the observed abundance of dark matter together with a smooth local distribution of dark matter \\cite{Cirelli-09}. Thus, an additional hypothesis is needed to boost the annihilation rate to the required levels and this must not change the present-day abundance of dark matter.  Such a boost is difficult to obtain from the effects of substructures in the local dark matter distribution. Recent numerical simulations predict this to be remarkably smooth \\cite{Vogelsberger-09}. A detailed analysis of the impact of substructure on the production of positrons by \\cite{Lavalle-08} came to a similar conclusion. The possibility of a nearby ``spike'' of dark matter produced by an intermediate mass black hole seems also a priori implausible \\cite{Bringmann-Lavalle-Salati-09}.  An alternative that has produced a plethora of papers in recent years is that of a Sommerfeld enhancement to the cross section produced by the mutual interaction of the annihilating dark matter particles. Such an interaction could be produced by a force carrier which might be any of the standard model weak force gauge bosons \\cite{Lattanzi-Silk-09} or a new force carrier \\cite{Arkani-Hamed-09}. For certain values of the parameters of these models, the enhancement is easily large enough to boost the cross section to the required values.  However, a large cross section has a significant impact on other observables and may violate other constraints. For instance, \\cite{Kamionkowski-Profumo-08} showed that for the case where the cross section increases as $1/v$ (a particular case of Sommerfeld-enhancement models), where $v$ is the relative velocity of the annihilating particles, there are severe constraints from measurements of the diffuse extragalactic gamma-ray background radiation and from CMB constraints on ionization and heating of the intergalactic medium (IGM) by annihilation in the first generation of halos. In this scenario, the boost factors required to fit the above anomalies would be inconsistent with current constraints. This problem is avoided, however, by more general cases of the Sommerfeld enhancement where the effect saturates at low velocities \\cite{Arkani-Hamed-09,Lattanzi-Silk-09}.  Recently, \\cite{Galli-09} and \\cite{Slatyer-Padmanabhan-Finkbeiner-09} analyzed constraints on the annihilation cross section from perturbations to the CMB angular power spectra resulting from heating and ionization of the photon-baryon plasma at recombination. They found interesting upper limits to models with Sommerfeld enhancement (see fig. 5 of \\cite{Galli-09}) that already rule out some extreme cases.  At higher redshifts, the effects of the Sommerfeld enhancement have scarcely been treated. For example,the annihilation of dark matter impacts the predictions from big-bang nucleosynthesis on the abundance of light elements (e.g. \\cite{Jedamzik-04}). This could put constraints on certain models with Sommerfeld enhancement. However, this has only been studied in passing \\cite{Hisano-09,Jedamzik-Pospelov-09}.  Also, the abundance of dark matter today is commonly assumed to be unaltered by this effect, apparently because the typical dark matter particle velocities at freeze-out are very large making the enhancement very close to one at that epoch \\cite{Arkani-Hamed-09,Kuhlen-Madau-Silk-09}. As we will show, this reasoning is flawed.  The thermodynamic equilibrium between matter and radiation in the early Universe would be perturbed by energy released during a certain process, dark matter annihilation for example. This equilibrium tends to be restored by different interaction mechanisms: Compton scattering, double Compton emission and bremsstrahlung radiation. The efficiency of these to fully restore equilibrium varies with redshift. For $z\\gtrsim2\\times10^6$ they are efficient enough to restore distortions in the energy spectrum and thus, the photon distribution is that of a black body with a slightly higher temperature than the one in the case of no energy injection. For lower redshifts these mechanisms can not restore the black body spectrum. In particular, for $5.1\\times10^4\\lesssim z \\lesssim2\\times10^6$, the spectrum is perturbed into a Bose-Einstein distribution with a chemical potential $\\mu$ \\cite{Illarionov-Siuniaev-75}.  In this paper, we revisit the impact of the Sommerfeld enhancement on the relic particle abundance and we show that its effect is not negligible. We also study, for the first time, $\\mu-$type distortions of the CMB spectrum due to energy deposition by dark matter annihilation in models with Sommerfeld enhancement.  The paper is organized as follows. In section 2 we summarize the Sommerfeld enhancement and describe how we include it in our calculations. The relic density calculation is set out in detail in section 3. In section 4, the $\\mu-$type distortion to the CMB from annihilation with Sommerfeld enhancement is calculated. Finally we present a summary and our conclusions in section 5. ",
        "conclusions": "The prospects for dark matter detection have increased considerably in recent years. There is a continual advance in the sensitivity of detectors on Earth that look for direct dark matter elastic scattering with nuclei, and in experiments that search indirectly for dark mater by looking for non-gravitational signatures of the byproducts of its hypothetical annihilation. Such technological improvements may lead, in the near future, to a definite proof of the existence of dark matter.  Perhaps dark matter has produced non-gravitational signals that have already been detected. The recently reported excesses of positrons in cosmic rays by PAMELA and of electrons+positrons by FERMI/ATIC/PPB-BETS can be explained by dark matter annihilation. Although other explanations with a different astrophysical origin are also possible, a solution based on dark matter is an attractive possibility.  However, such solution seems to require an attractive force between the dark matter particles to enhance their annihilation through the Sommerfeld mechanism. For a Yukawa interaction via a single scalar, the magnitude of the enhancement depends on the coupling strength  of the interaction $\\alpha_c$, on the mass ratio of the force carrier to the dark matter particle $m_{\\phi}/m_{\\chi}$ and on the relative velocities of the annihilation particles $\\beta$.  Several recent papers have computed the boost to the annihilation due to a Sommerfeld enhancement (in the Galactic halo and/or in its subhalos) \\cite{Lattanzi-Silk-09,Arkani-Hamed-09,Bovy-09}. One of the main aims of these works is to show that for certain values of $\\alpha_c$ and $m_{\\phi}/m_{\\chi}$, this mechanism is able to produce large enough boosts to explain the cosmic ray anomalies.  Only a handful of studies have addressed in detail the impact that such an enhancement has in the early Universe. In \\cite{Kamionkowski-Profumo-08}, the authors found that if the cross section increases with decreasing relative velocity as $1/\\beta$ (which is valid in a certain regime for general Sommerfeld enhancement models), dark matter annihilation in the first halos would heat and ionize the IGM, violating current constraints from the CMB. As was later pointed out by \\cite{Lattanzi-Silk-09} and \\cite{Arkani-Hamed-09}, this problem is alleviated in more general models where the enhancement saturates at low velocities. In \\cite{Galli-09} and \\cite{Slatyer-Padmanabhan-Finkbeiner-09}, a constraint on the annihilation cross section was obtained by considering limits on the energy deposition by annihilation at recombination. The constraint reported by \\cite{Slatyer-Padmanabhan-Finkbeiner-09} is $\\langle\\sigma v\\rangle_{REC}<3.6\\times10^{-24}{\\rm cm}^3{\\rm s}^{-1}(m_{\\chi}/1{\\rm TeV})/f_{REC}$, where $f_{REC}$ is an average efficiency of energy injection into the IGM by annihilation at recombination.  In the present paper, we have analyzed the impact of dark matter annihilation with Sommerfeld enhancement at higher redshift.  In the first place, we have revisited the calculation of the dark matter particle abundance by solving the Boltzmann equation from freeze-out through the epoch of kinetic decoupling, including a full solution to the Schr\\\"odinger equation. Contrary to previous claims \\cite{Arkani-Hamed-09,Kuhlen-Madau-Silk-09}, we have found a significant suppression of the relic density, in agreement with a recent work by \\cite{Dent-Dutta-Scherrer-09}. This suppression is particularly important near resonances, which are typically invoked to explain the cosmic ray anomalies.  We found that to fit the observed dark matter abundance, the normalization of the cross section needs to be lowered by up to a factor of 10 compared to the case without enhancement (see Fig.~\\ref{cross_section}). The result depends on the coupling strength and the proximity to a resonance. Exploring a broad range of dark matter particle masses and kinetic decoupling temperatures, we found a minor-to-medium impact of these parameters on our results; variations on $T_{kd}$ have the strongest impact.  Secondly, we have calculated the amount of energy deposited by dark matter annihilation into the radiation plasma in the redshift range $5.1\\times10^4<z<2\\times 10^6$. Energy injection at this epoch would create a distortion in the CMB energy spectrum first pointed out by \\cite{Illarionov-Siuniaev-75}. The COBE/FIRAS experiment has put constraints on this Bose-Einstein $\\mu-$type distortion. The upper limit on the chemical potential associated to the distortion is $\\vert\\mu\\vert\\leq9.0\\times10^{-5}$ at the $2\\sigma$ level \\cite{Fixsen-96}. We have found that very near resonances, annihilation with Sommerfeld enhancement is already ruled out by this constraint on $\\mu$ (see Fig~\\ref{mu}). Quantitatively, the ``safe'' range of proximity to a resonance depends on the particular model, higher values of dark matter mass and kinetic decoupling temperature allow a closer proximity to the resonance (see Fig.~\\ref{mu_resonance}). In the case where the force carrier is massless, the Yukawa interaction reduces to a Coulomb one. Our findings are much more stringent in this case with values of $\\alpha_c>6\\times10^{-2}$ already ruled out (see Fig.~\\ref{mu_coulomb}).  Improved upper limits on the $\\mu$ parameter by a null distortion detection in the CMB spectrum would rule out a larger region of the parameter space for the Yukawa interaction. In \\cite{Fixsen-Mather-02}, it was pointed out that an improvement close to two orders of magnitude has already been possible for a number of years, and \\cite{Mather-07} suggests that another order of magnitude is perhaps within reach. An upper limit on $\\vert\\mu\\vert$ of the order of $10^{-7}$ would certainly exclude large regions of the parameter space. It would exclude to a large extent near-resonance regions, which are the ones that produce the largest boosts to the annihilation. On the other hand, the detection of a distortion could possibly tell us something about the parameters of the Yukawa interaction and the nature of the force carrier $\\phi$.  In summary, our results indicate that for a given set of parameters $(m_{\\chi},T_{kd},m_{\\phi},\\alpha_c)$, it is necessary to compute in detail the relic density to obtain the range of values for the normalization to the cross section that are compatible with current estimates of the dark matter abundance. Once this normalization is known, the energy input producing a $\\mu-$type distortion in the CMB spectrum can be computed to check whether the particular model violates current constraints. Only for allowed models can a boost factor for local dark matter annihilation be computed and advocated.  Our findings show that the local boosts reported in the literature need to be renormalized to the proper value of $\\langle\\sigma v\\rangle_S$ implied by the observed relic density. This renormalization can exceed a factor of 10 in extreme cases.  \\begin{figure} \\centering \\includegraphics[height=8.5cm,width=10cm]{./fig6.ps} \\caption{The values of the boost factor, relative to $\\langle\\sigma v\\rangle_S=3\\times10^{-26}{\\rm cm}^3{\\rm s}^{-1}$ for the parameter scan of the Yukawa interaction. These values are obtained according to the cross section values of Fig.~\\ref{cross_section}, which are consistent with $\\Omega_{\\chi}h^2=0.1143$. A Maxwell-Boltzmann velocity distribution with $\\sigma_v=5\\times10^{-4}$ for a neutralino with $m_{\\chi}=100$GeV and $T_{kd}=8\\times10^{-3}$GeV was used. The values are color-coded logarithmically according to the scale on the right.} \\label{boost_factor} \\end{figure}  A recent analysis of the PAMELA and FERMI data by \\cite{Bergstrom-Edsjo-Zaharijas-09} suggests that boost factors, over a standard value of $\\langle\\sigma v\\rangle_S=3\\times10^{-26}{\\rm cm}^3{\\rm s}^{-1}$, of order 1000 or higher are required to explain both datasets simultaneously. We note that smaller boost factors ($\\sim 100$) can still fit the PAMELA data, but not the FERMI one, for a 100~GeV neutralino, which is the standard case we have considered here. These large boosts are also found in the analysis by \\cite{Meade-09}. It has been argued in the past that these boosts can be achieved only by invoking Sommerfeld-enhanced annihilation. For instance, \\cite{Arkani-Hamed-09} and \\cite{Bovy-09} obtain maximum boost factors $\\sim1000$ by assuming a local Maxwell-Boltzmann velocity distribution with a velocity dispersion of $150~{\\rm kms}^{-1}$ ($5\\times10^{-4}c$), which roughly corresponds to the estimated local dark matter velocity dispersion. We find that these boosts are modified once the proper cross section is used.  In Fig.~\\ref{boost_factor} we show the boost factors, i.e., the multiplicative factor to $\\langle\\sigma v\\rangle_S=3\\times10^{-26}{\\rm cm}^3{\\rm s}^{-1}$, that we find for a Maxwell-Boltzmann velocity distribution with $\\sigma_v=5\\times10^{-4}$ for the parameter space of the Yukawa interaction. The figure is for $m_{\\chi}=100$GeV, $T_{kd}=8\\times10^{-3}$GeV and is consistent with a relic abundance $\\Omega_{\\chi}h^2=0.1143$. Keep in mind, however, that changes in the neutralino mass have almost no impact on the normalization of the cross section. Different kinetic decoupling temperatures and variations on $\\Omega_{\\chi}h^2$ have also only a small effect ($<60\\%$ in combination). Thus, Fig.~\\ref{boost_factor} is approximately correct for all the cases considered in this paper. The figure shows that even in extreme cases, for resonances in the upper right of the figure, the boost factors are $<500$. In more favored regions of the parameter space $\\alpha_c\\gtrsim10^{-2}$, $m_{\\phi}/m_{\\chi}\\lesssim10^{-3}$ (see for example \\cite{Bovy-09}), the boost factors are $\\lesssim100$.  This result suggests that additional assumptions are needed to account for the boost needed to explain the cosmic ray anomalies by dark matter annihilation alone. The inclusion of colder substructures to the overall smooth component with higher densities and lower velocity dispersions and thus, higher boost factors, could perhaps solve the issue. However, as found in recent high-resolution N-body simulations, the local dark matter distribution is rather smooth \\cite{Vogelsberger-09}. Typical estimates on such an additional boost factor due to substructure in the galactic halo (albeit without Sommerfeld enhancement) are of the order of 1.4, unless a subhalo happens to be very close to Earth, in which case this boost could be larger than 10 \\cite{Diemand-08}. These estimates do not include a possible further amplification due to the scaling of the Sommerfeld enhancement with velocity dispersion. However, the constraints found by \\cite{Slatyer-Padmanabhan-Finkbeiner-09} based on energy deposition at recombination suggest that the enhancement must be already close to saturation for $\\sigma_v=150{\\rm kms}^{-1}$, thus such a further amplification seems implausible. Therefore, is not clear that the inclusion of colder substructures can account for the additional boost.  Finally, we mention that we have created a web application that solves the relevant equations described in this paper. It allows the user to compute individual Sommerfeld boosts, thermally averaged enhancements, cross section and $\\mu$ values, among other quantities. The interested reader can find this at http://www.mpa-garching.mpg.de/$\\sim$vogelsma/sommerfeld/"
    },
    "0910/0910.2681.txt": {
        "abstract": "We explore how the local environment is related to properties of active galactic nuclei (AGNs) of various luminosities.  Recent simulations and observations are converging on the view that the extreme luminosity of quasars, the brightest of AGNs, is fueled in major mergers of gas-rich galaxies.  In such a picture, quasars, the highest luminosity AGNs, are expected to be located in regions with a higher density of galaxies on small scales where mergers are more likely to take place.  However, in this picture, the activity observed in low-luminosity AGNs is due to secular processes that are less dependent on the local galaxy density.  To test this hypothesis, we compare the local photometric galaxy density on kiloparsec scales around spectroscopic type I and type II quasars to the local density around lower-luminosity spectroscopic type I and type II AGNs.  To minimize projection effects and evolution in the photometric galaxy sample we use to characterize AGN environments, we place our random control sample at the same redshift as our AGNs and impose a narrow redshift window around both the AGNs and control targets.  Our results support these merger models for bright AGN origins.  We find that the brightest sources have overdensities that increase on the smallest scales compared to dimmer sources.  In addition, we investigate the nature of the quasar and AGN environments themselves and find that the increased overdensity of early-type galaxies in the environments of bright type I sources suggests that they are located in richer cluster environments than dim sources.  We measure increased environment overdensity with increased quasar black hole mass, consistent with the well-known $M_{\\rm DMH}-\\MBH$ relationship, and find evidence for quenching in the environments of high accretion efficiency type I quasars. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.0632_arXiv.txt": {
        "abstract": "{% The Square Kilometre Array (SKA) is intended as the next-generation radio telescope and will address fundamental questions in astrophysics, physics, and astrobiology.  The international science community has developed a set of Key Science Programs: (1)~Emerging from the Dark Ages and the Epoch of Reionization, (2)~Galaxy Evolution, Cosmology, and Dark Energy, (3)~The Origin and Evolution of Cosmic Magnetism, (4)~Strong Field Tests of Gravity Using Pulsars and Black Holes, and (5)~The Cradle of Life/Astrobiology. In addition, there is a design philosophy of ``exploration of the unknown,'' in which the objective is to keep the design as flexible as possible to allow for future discoveries.  Both a significant challenge and opportunity for the \\hbox{SKA} is to obtain a significantly wider field of view than has been obtained with radio telescopes traditionally.  Given the breadth of coverage of cosmic magnetism and galaxy evolution in this conference, I highlight some of the opportunities that an expanded field of view will present for other Key Science Programs.}  \\FullConference{Panoramic Radio Astronomy: Wide-field 1-2 GHz research on galaxy evolution - PRA2009\\\\ June 02 - 05 2009\\\\ Groningen, the Netherlands}  \\begin{document} ",
        "introduction": "\\label{sec:jl.intro}  In the $20^{\\mathrm{th}}$ Century, we discovered our place in the Universe.  We learned that it was much bigger than we imagined and much more exotic.  Beyond our Milky Way, the Universe is filled with galaxies---each their own island universe.  They range in size from dwarf galaxies barely able to survive near their larger neighbors to giant elliptical galaxies, orders of magnitudes larger than the Milky Way.  These galaxies of stars also contain a multitude of other components including gas with a wide range of temperatures; compact objects including white dwarfs, neutron stars, and black holes; and planets.  Over the course of the century, black holes moved from a theoretical curiosity to a well-recognized endpoint of stellar evolution and a likely fundamental component of the centers of galaxies, with the potential to power immense jets of relativistic particles that affect their surroundings.  By the end of the century, we were beginning to unveil the basic structure and processes of the Universe in which these objects are embedded, including evidence of its origin and the still mysterious properties of dark matter and dark energy.  Our probes of the Universe have expanded dramatically as well. Electromagnetic radiation has been detected from celestial objects at frequencies below~1~MHz ($\\lambda \\sim 300$~m) to energies exceeding 1~\\hbox{TeV}.  Moreover, the range of possible signals has expanded beyond just electromagnetic radiation.  Cosmic rays rain down on the Earth, some with energies approaching those of macroscopic objects. Gravitational radiation has been detected indirectly, and numerous potential classes of sources have been suggested, with the expectation that the Earth is awash in gravitational waves.  Neutrinos have been detected from both the Sun and supernova \\hbox{1987A}, and many of the processes that generate high energy cosmic rays should also produce a spectrum of high-energy neutrinos.  In the $21^{\\mathrm{st}}$ Century, we seek to understand the Universe we inhabit.  To do so will require a suite of powerful new instruments, on the ground and in space, operating across the entire electromagnetic spectrum and for multiple decades.  Observations at centimeter- to meter wavelengths have provided deep insight to a wide range of phenomena ranging from the solar system to the most distant observable celestial emission. This long and rich record of important discoveries in the radio spectrum, including 3 Nobel prizes, has been possible since many of the relevant physical phenomena can only be observed, or understood best, at these wavelengths.  These phenomena include the cosmic microwave background (CMB), quasars, pulsars, gravitational waves, astrophysical masers, magnetism from planets through galaxies, the ubiquitous jets from black holes and other objects, and the spatial distribution of hydrogen gas, the predominant baryonic constituent of the Universe.  Moreover, through the invention of aperture synthesis, also recognized by the Nobel committee, radio astronomy has reached unprecedented levels of imaging resolution and astrometric precision, providing the fuel for further discovery.  With only a handful of exceptions, radio telescopes and arrays have been limited to apertures of about~$10^4$~m${}^2$, constraining, for instance, studies of the 21-centimeter hydrogen emission to the nearby Universe ($z \\sim 0.2$)~\\cite{cr04}.  Contemporaneous with the astronomical discoveries in the latter half of the $20^{\\mathrm{th}}$ Century have been technological developments that offer a path to substantial improvements in future radio astronomical measurements. Among the improvements are mass production of centimeter-wavelength antennas enabling apertures potentially 100 times larger than previously available, fiber optics for the transmission of large volumes of data, high-speed digital signal processing hardware for the acquisition and analysis of the signals, and computational improvements leading to massive processing and storage.  These new technologies, combined with dramatically improved survey speeds and the other advances, can open up an enormous expanded volume of discovery space, providing access to many new celestial phenomena and structures, including 3-dimensional mapping of the web of hydrogen gas through much of cosmic history ($z \\sim 2$).  The realization that radio astronomy was on the doorstep of a revolutionary age of scientific breakthrough has led the international community to investigate this opportunity in great detail over the last decade.  That coordinated effort, involving a significant fraction of the world's radio astronomers and engineers, has resulted in the Square Kilometre Array (SKA) Program (Figure~\\ref{fig:ska}), an international roadmap for the future of radio astronomy over the next two decades and one for which access to a wide field of view is an integral part of the science.  \\begin{figure} \\begin{center} \\includegraphics[width=0.9\\textwidth]{SKA_cores_2009.eps} \\end{center} \\vspace*{-3ex} \\caption{An artist's impression of the core of the SKA illustrating the various technologies over the frequency range 70~MHz to~10~GHz. All of these technologies would enable various levels of wide-field imaging.} \\label{fig:ska} \\end{figure}  From its inception, development of the SKA Program has been a global endeavor.  In the early 1990s, there were multiple, independent suggestions for a ``large hydrogen telescope.''  It was recognized that probing the fundamental baryonic component of the Universe much beyond the local Universe would require a substantial increase in collecting area.  The IAU established a working group in~1993 to begin a worldwide study of the next generation radio observatory.  Since that time, the effort has grown to comprise 19 countries and more than 50 institutes, including about~200 scientists and engineers. ",
        "conclusions": ""
    },
    "0910/0910.0404_arXiv.txt": {
        "abstract": "The kinematics of the extra-planar neutral and ionised gas in disc galaxies shows a systematic decline of the rotational velocity with height from the plane (vertical gradient).  This feature is not expected for a barotropic gas, whilst it is well reproduced by baroclinic fluid homogeneous models.  The problem with the latter is that they require gas temperatures (above $10^5$ K) much higher than the temperatures of the cold and warm components of the extra-planar gas layer.  In this paper, we attempt to overcome this problem by describing the extra-planar gas as a system of gas clouds obeying the Jeans equations. In particular, we consider models having the observed extra-planar gas distribution and gravitational potential of the disc galaxy NGC\\,891: for each model we construct pseudo-data cubes and we compare them with the HI data cube of NGC\\,891. In all cases the rotational velocity gradients are in qualitative agreement with the observations, but the synthetic and the observed data cubes of NGC\\,891 show systematic differences that cannot be accommodated by any of the explored models.  We conclude that the extra-planar gas in disc galaxies cannot be satisfactorily described by a stationary Jeans-like system of gas clouds. ",
        "introduction": "Observations of several spiral galaxies at various wavelengths have revealed massive gaseous haloes surrounding the galactic discs.  This extra-planar gas is multiphase: it is detected in \\ion{H}{I} (e.g. Swaters, Sancisi \\& van der Hulst 1997), in H$\\alpha$ (Rand 2000; Rossa et al. 2004) and in X-ray observations (Wang et al. 2001; Strickland et al. 2004). Thanks to high-sensitivity \\ion{H}{I} observations of some edge-on galaxies, like NGC 891 (Oosterloo, Fraternali \\& Sancisi 2007), we can trace these haloes up to $10$ - $20$ kpc from the plane and we can study their kinematics in detail. A remarkable feature of the extra-planar gas is its regularly decreasing rotational velocity $u_{\\varphi}$ at increasing distances from the galactic plane (vertical gradient). Fraternali et al.\\ (2005) found that in NGC 891 the vertical gradient is $\\sim\\, -15 {\\rm km \\,s} ^{-1} {\\rm kpc}^{-1}$ (see also Heald, Rand \\& Benjamin et al. 2007). For other galaxies (e.g. NGC 2403) gradients of the same order have been estimated (see Fraternali 2008).  The dynamical state of the extra-planar gas can provide useful insights to understand its origin.  Two main frameworks have been proposed: the galactic fountain (Shapiro \\& Field 1976; Bregman 1980) and the cosmological accretion (Oort 1970; Binney 2005).  In the galactic fountain model, partially ionized gas is ejected from the disc by supernova explosions and stellar winds, it travels through the halo and eventually falls back to the disc. Due to difficulties in running high-resolution hydrodynamical simulations of the whole galactic disc (see e.g.\\ Melioli et al.\\ 2008), the gas in the galactic fountain model is often treated ballistically, i.e. as a collection of non-interacting clouds subject to the gravitational field of the galaxy (Collins, Benjamin \\& Rand 2002). The ballistic galactic fountain is able to reproduce the gas distribution observed in spiral galaxies, but in general the predicted vertical gradient of $u_{\\varphi}$ is significantly lower than observed (Fraternali \\& Binney 2006).  In the accretion models, the extra-planar gas is the result \\rev{of the infall of cool intergalactic gas}. The two scenarios are not mutually exclusive: in fact it has been argued that the contribution of accretion is limited to only $10-20\\%$ of the total mass of extra-planar gas (Fraternali \\& Binney 2008).  In a complementary approach, some authors (Benjamin 2002; Barnab\\`{e} et al. 2006), have focused on the construction of models for the extra-planar gas under the assumption of stationarity. In this approach the problem of the origin of the extra-planar gas is not addressed, and the main effort is to clarify the global dynamical state of the gas. Usually, in the stationary approach, the extra-planar gas is described as an homogeneous fluid in permanent rotation. In the simplest of fluid homogeneous models, vertical and radial motions of the gas are neglected and therefore the vertical gravitational field of the galaxy is balanced only by the pressure gradient of the gas. In other studies, turbulence has also been added as an extra-pressure term (e.g., Koyama \\& Ostriker 2008). According to the Poincar\\'{e}-Wavre theorem (Lebovitz 1967; Tassoul 1978), a vertical gradient in the rotational velocity can be present only if the pressure of the medium is not stratified on the density, therefore in fluid homogeneous models the gas distribution is necessarily \\textit{baroclinic} (Barnab\\`{e} et al. 2006; see also Waxman 1978). Remarkably, in their study of NGC 891, Barnab\\`{e} et al. (2006) found that if the baroclinic configuration is built from general physical arguments, the model rotational velocity well reproduces the observed rotation curves at different heights over the galactic disc and for a large radial range. However, the temperature predicted for the system is $> 10^{5}$ K, well above that of the neutral gas.  Here we investigate if and how this ``temperature problem'' can be solved. We make use of the fact that stationary fluid equations for a gaseous system in permanent rotation are, from a formal point of view, identical to the stationary Jeans equations for an axisymmetric system with isotropic velocity dispersion. Thus, also in this case, a decrease of rotational velocities with increasing distances from the mid-plane is obtained when the density field is built using the approach described by Barnab\\`{e} et al. (2006). In this paper we analyse this alternative interpretation of the baroclinic fluid homogeneous models. In practice we consider a ``gas'' of cold \\ion{H}{I} clouds described by the stationary Jeans equations, where the thermal pressure, needed for the vertical balance of the galaxy gravitational field, is replaced by the velocity dispersion of the clouds.  \\rev{ In the present approach, some preliminary consideration on the physical state of the clouds is important. In fact, \\ion{H}{I} halo clouds are often assumed to be embedded in, and in pressure equilibrium with, a hot medium (corona) that is about a thousand times less dense and provides the pressure required to confine them (Spitzer 1956; Wolfire et al. 1995). In the Milky Way, there is plenty of evidence for the existence of this medium, such as emission from highly ionised metals (e.g. Sembach et al.\\ 2003) and the head-tail morphology of individual High Velocity Clouds (HVCs, e.g. Br{\\\"u}ns et al. 2000). Due to the high density contrast, the pressure-confined clouds cannot be supported by the pressure of the external medium and must be moving more or less ballistically in a way akin to water drops in the air (Bregman 1980). This assumption is commonly employed in galactic fountain models (e.g. Collins et al. 2002). On the time-scale taken for a cloud to move about 1000 times its length, the trajectory may deviate significantly from the ballistic one due to the interactions with the external medium. This time-scale depends on the cloud mass but it is longer than a dynamical time for the larger clouds (Fraternali \\& Binney 2008). The typical mass of an halo cloud is difficult to estimate. In the Milky Way the new determinations of distances for the HVCs give masses from a few $\\times 10^4 \\mo$ to $10^6 \\mo$ (e.g. Wakker et al. 2008), for external galaxies the mass resolution is often above $10^6 \\mo$ but in M31 several clouds with masses down to $\\sim 1 \\times 10^5 \\mo$ have been observed (Thilker et al. 2004). We estimate that for a cloud of $10^5 \\mo$, the drag time scale is more than 10 times the dynamical time and it is therefore fair to treat the system as purely dynamical. Moreover, if the \\ion{H}{I} halo of a galaxy like NGC\\,891 is made up of clouds with this mass and radii of $100$ pc, the collision time between clouds turns out to be about 5 times the dynamical time and the system can be considered collision-less. }  The substitution of a fluid system with a gas of clouds is also not trivial from the point of view of the comparison with the observations, as several delicate issues arise when considering \\ion{H}{I} observations (see Sect. \\ref{constr}). To overcome these problems we follow Fraternali \\& Binney (2006, 2008) and construct, as an output for each of the models investigated, a \\textit{pseudo-data cube} with the same resolution and total flux as the \\textit{data cube} of the \\ion{H}{I} observations. This procedure assures a full control of projection and resolution effects. Moreover, the comparison with the raw data removes all the intermediate stages of data analysis and the associated uncertainties. For completeness, we also investigate some phenomenological anisotropic models.  The paper is organized as follows: in Section 2 we illustrate the Jeans-based interpretation of the baroclinic solutions, introducing the anisotropic case which has no analogous fluid counterpart, and we discuss the conditions required for these solutions to have a negative vertical velocity gradient. In Section~3 we present the method adopted to construct isotropic and anisotropic models for the edge-on galaxy NGC 891, whilst in Section~4 we compare their predictions with the \\ion{H}{I} observations. Section~5 is devoted to the discussion of the results, and Section~6 concludes. ",
        "conclusions": "In this paper, motivated by the results of Barnab\\`e et al. (2006), we investigated the possibility that the extra-planar gas in spiral galaxies can be modelled as a ``gas'' composed by cold \\ion{H}{I} clouds that follows the stationary Jeans equations. \\rev{In doing this we are assuming that the pressure-bound clouds are moving almost ballistically in the halo. This assumption is valid in the limit of massive clouds having negligible rates of mass exchange with the corona. }  In this alternative interpretation of fluid baroclinic models the thermal pressure is replaced by an isotropic velocity dispersion tensor, so that the problem of high temperature of gaseous homogeneous models is eliminated. We have also extended the discussion to simple phenomenological models with anisotropic velocity dispersion tensors. We constructed both isotropic and anisotropic models in the well-constrained gravitational field of the spiral galaxy NGC 891, and compared their predictions to the observed kinematics of the extra-planar gas in that galaxy.  For each model we built a pseudo-data cube with the same resolution and total flux as the observations.  The main results of our analysis can be summarized as follows:  \\begin{enumerate}  \\item The adopted functional form of the cloud density distribution, taken by the observations, leads to physically acceptable solutions in all the models investigated. The cloud density distribution is centrally depressed, and, in order to match the vertical extension of the \\ion{H}{I} halo of NGC 891, it has higher scale-height than that used by Barnab\\`{e} et al. (2006).  \\item All the models computed show a negative vertical gradient in the rotational velocity, the distinctive feature of the kinematics of the extra-planar gas.  \\item The support against the vertical gravitational field of the galaxy requires a $\\sigma_{z} \\simeq 50 - 100\\, \\rm{km\\, s^{-1}}$ and therefore the line-of-sight velocity dispersion of the isotropic model is a factor $\\sim\\,3-4$ higher than that observed.  \\item With the introduction of the anisotropy it is possible to restrict the line-of-sight velocity dispersion to the observed values, but the predicted vertical gradient in the rotational velocity is somewhat too shallow and other features of the data cube are not fully reproduced.  \\end{enumerate}  We conclude that the dynamics of extra-planar gas in a galaxy like NGC\\,891 is not fully described by any of the stationary models considered here.  However a model with an anisotropic velocity dispersion tensor, which mimics a galactic fountain is the preferable among all. \\rev{The fact that none of the stationary Jeans models analysed here can reproduce all the features of the observed (extra-planar) gas kinematics might suggest that the cloud motion is not purely ballistic, and that the interaction between the clouds and the coronal gas, perhaps in the form of mass exchange, plays an important dynamical role.}"
    },
    "0910/0910.4011_arXiv.txt": {
        "abstract": "We present $25''$ resolution radio images of five Lynds Dark Nebulae (L675, L944, L1103, L1111 and L1246) at 16\\,GHz made with the Arcminute Microkelvin Imager (AMI) Large Array. These objects were previously observed with the AMI Small Array to have an excess of emission at microwave frequencies relative to lower frequency radio data. In L675 we find a flat spectrum compact radio counterpart to the 850\\,$\\mu$m emission seen with SCUBA and suggest that it is cm-wave emission from a previously unknown deeply embedded young protostar. In the case of L1246 the cm-wave emission is spatially correlated with 8\\,$\\mu$m emission seen with \\emph{Spitzer}. Since the MIR emission is present only in \\emph{Spitzer} band 4 we suggest that it arises from a population of PAH molecules, which also give rise to the cm-wave emission through spinning dust emission. ",
        "introduction": "The complete characterization of microwave emission from spinning dust grains is a key question in both astrophysics and cosmology. It probes a region of the electromagnetic spectrum where a number of different astrophysical disciplines overlap. It is important for CMB observations in order to correctly characterise the contaminating foreground emission; for star and planetary formation it is important because it potentially probes a regime of grain sizes that is not otherwise easily observable.  Although a number of objects have now been found to exhibit anomalous microwave emission, attributed to spinning dust, it is still unclear what differentiates those objects from the many other seemingly similar targets that do not show the excess. In the specific case of dark clouds the recent AMI sample (AMI Consortium: Scaife et~al.\\ 2009; hereinafter Paper I) of fourteen Lynds Dark Nebulae found an excess in only five.  It has been suggested that cm-wave emission from spinning dust is emitted by a population of ultra-small grains (Draine \\& Lazarian 1998). These ultra-small grains are thought to exist mainly in the form of single polycyclic aromatic hydrocarbon (PAH) molecules. PAH molecules are generally detected through their narrow line emission features in the MIR. For these emission features to be observed the PAH molecules must be exposed to a strong source of UV flux. Since this flux is generally absent in the case of dark clouds, the microwave emission from the rotation of PAH molecules may be the only way to study the very small grain population in these objects.  It is also known that radio continuum emission in dark clouds may arise from ionized gas associated with a stellar outflow. When a luminous star is present this arises either as the result of a compact {\\sc{Hii}} region or an ionized stellar wind. In the case of very young low luminosity stars radio continuum emission may be also be detected. In this instance it is generally attributed to the presence of a partially ionized ($0.02\\leq x_{\\rm e} \\leq 0.35$; Bacciotti \\& Eisl{\\\"o}ffel 1999) stellar wind (Wright \\& Barlow 1975; Panagia \\& Felli 1975), or possibly a neutral wind which has been shock-ionized further from the central source by impacting on a dense obstacle (Curiel et~al.\\ 1989).  In this paper we present follow-up observations of the five AMI Small Array (SA) spinning dust detections (Paper I) at higher resolution with the AMI Large Array (LA) over the same frequency range. All co-ordinates in this paper are J2000.0. ",
        "conclusions": "L675 and L1246 have archival \\emph{Spitzer} IRAC data, which shows in both cases a significant amount of emission in Band 4 (6.4--9.4\\,$\\mu$m) and very little in the other three (3.2--3.9, 4.0--5.0 and 4.9--6.4\\,$\\mu$m, respectively). In the case of L675 this emission is present on a very large scale, see Fig.~\\ref{fig:l675spit}. The emission seen at 16\\,GHz with the AMI SA appears on a similar scale, however the small field of view of the \\emph{Spitzer} data precludes a more detailed comparison. L1246 shows an arc of emission at 16\\,GHz which is also evident in \\emph{Spitzer} IRAC Band 4, see Fig.~\\ref{fig:l1246spit}. This emission is again not present in Bands 1--3. In Band 4 it is present as an arc, coincident with that seen at 16\\,GHz in the AMI LA data.  \\emph{Spitzer} Band 4 contains two of the PAH emission lines, including the strongest (7.7\\,$\\mu$m). Of the three other \\emph{Spitzer} bands only Band 1 contains an emission line (3.3\\,$\\mu$m) and for ionized PAHs this line is expected to be significantly weaker. It is probable therefore that the MIR correlated cm-wave data seen in the AMI maps is a consequence of spinning dust emission from a population of ionized PAH molecules. Neutral PAH molecules do not in general possess a permanent dipole moment and are therefore not expected to have rotational emission (Tielens 2008). This emission, the mechanism of which is described in detail by Draine \\& Lazarian (1998), arises from the intrinsic dipole moments of small dust grains, most likely to be PAH molecules, which emit power when they rotate. This rotation has a variety of contributing factors, the relative importance of which varies with grain environment. However, in the majority of cases excitation through collision with ions is predominant.  In the case of L675A, we must consider the possibility that we are observing a coincidental extragalactic radio source. Using the extended 9C survey 15\\,GHz source counts (Waldram et~al.\\ 2009), where $n(S) = 51 (S/{\\rm{Jy}})^{-2.15}$\\,Jy$^{-1}$\\,sr$^{-1}$, the probability that a source with flux density greater than 2\\,mJy lies within the FWHM of the AMI LA primary beam is $0.12$, and only 0.01 within the SCUBA field. It is likely therefore that the radio source L675A is associated with the SCUBA core.  A further question is whether the cm-wave emission might be explained by thermal (Planckian) dust emission. A single greybody spectrum with a dust temperature, $T_{\\rm{d}} \\approx 27$\\,K, might be used to explain the LA flux density, however it would require a $\\beta$ of 0.6. Such a value would be unusual even for objects known to possess flattened dust tails, such as protoplanetary disks. This simple fit also neglects the flux lost by the AMI LA baseline distribution. SA observations have already shown this source to possess a significant amount of extended emission which would make this scenario even more unlikely.  The presence of a neutral or partially ionized wind from an outflow source that has been shocked through encountering a dense obstacle (Torrelles et~al.\\ 1985; Rodr{\\'i}guez et~al.\\ 1986) is used to understand the spectral indices seen towards exciting sources in the radio regime (Curiel et~al.\\ 1990; Cabrit \\& Bertout 1992). This model allows a spectral index range of $0.1$ (optically thin) to -2 (optically thick), which explains results which deviate from the value of $\\alpha = -0.6$ required by a spherically symmetric ionized wind (Wright \\& Barlow 1975; Panagia \\& Felli 1975). Using this model as described in Curiel et~al.\\ (1989; 1990) the radio emission is expected to be optically thin ($\\tau=0.1$), consistent with the spectral index seen across the AMI band. Assuming a distance of 300\\,pc and a stellar wind with a wind speed of 200\\,km\\,s$^{-1}$, we can calculate that the AMI flux densities towards L675A are consistent with a mass loss of $3.5\\times 10^{-7}$\\,M$_{\\odot}$\\,yr$^{-1}$. A mass loss such as this implies a mechanical luminosity from the wind of $L_{\\rm{mech}} \\approx 1.1$\\,L$_{\\odot}$, comparable to the values found by Curiel et~al for L1448.  The nature of the emission seen towards L944 with the AMI LA is uncertain. The spectral index of this emission is consistent with spinning dust emission or alternatively the optically thick component of free--free spectrum. Such a free--free spectrum might be exhibited at 16\\,GHz by ultra-compact {\\sc Hii} regions. However a turn-over frequency above 16\\,GHz would have an extremely high mass and should therefore be obvious in sub-mm observations. This needs to be confirmed by either higher radio frequency measurements in order to measure the optically thin region of the spectrum and the turn-over, or sub-mm measurements to place constraints on the mass of such a region.  In conclusion, we have used the AMI LA to observe a sample of five Lynds Dark Nebulae selected as candidates for spinning dust emission from the AMI SA sample of Lynds Dark Nebulae (Paper I). Towards two of these clouds (L1103 and L1111) we detect only patchy diffuse emission characteristic of the presence of a larger structure which has been mostly resolved out.  Towards L675 we have observed flat spectrum compact cm-wave emission coincident with the SCUBA 850\\,$\\mu$m emission from the same region. These characteristics suggest that this source is associated with a stellar wind from a deeply embedded young protostar.  We detect extended cm-wave emission to the North of the L944 SMM-1 protostar which displays spectral behaviour consistent with either spinning dust, or alternatively a collection of ultracompact {\\sc Hii} regions.  L1246 shows an arc of cm-wave emission which is coincident with emission seen in \\emph{Spitzer} Band 4. We suggest that this is an example of emission from a population of PAH molecules, seen in emission lines in the \\emph{Spitzer} data, and emission as a consequence of rapid rotation of the molecules in the cm-wave data."
    },
    "0910/0910.1919_arXiv.txt": {
        "abstract": "Spectroscopic long-slit observations of the dwarf Irr~galaxy IC~10 were conducted at the 6-m Special Astrophysical Observatory telescope with the SCORPIO focal reducer. The ionized-gas emission spectra in the regions of intense current star formation were obtained for a large number of regions in IC~10. The relative abundances of oxygen, N$^+$, and S$^+$ in about twenty HII~regions and in the synchrotron superbubble were estimated. We found that the galaxy-averaged oxygen abundance is $12+\\log(\\textrm{O/H}) = 8.17 \\pm 0.35$ and the metallicity is $Z=0.18 \\pm 0.14 Z_{\\odot} $. Our abundances estimated from the strong emission lines are found to be more reliable than those obtained by comparing diagnostic diagrams with photoionization models. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.1634_arXiv.txt": {
        "abstract": "The accretion process onto spinning objects in Kerr spacetimes is studied with numerical simulations. Our results show that accretion onto compact objects with Kerr parameter (characterizing the spin) $|a| < M$ and $|a| > M$ is very different. In the super-spinning case, for $|a|$ moderately larger than $M$, the accretion onto the central object is extremely suppressed due to a repulsive force at short distance. The accreting matter cannot reach the central object, but instead is accumulated around it, forming a high density cloud that continues to grow. The radiation emitted in the accretion process will be harder and more intense than the one coming from standard black holes; e.g. $\\gamma$-rays could be produced as seen in some observations. Gravitational collapse of this cloud might even give rise to violent bursts. As $|a|$ increases, a larger amount of accreting matter reaches the central object and the growth of the cloud becomes less efficient. Our simulations find that a quasi-steady state of the accretion process exists for $|a|/M \\gtrsim 1.4$, independently of the mass accretion rate at large radii. For such high values of the Kerr parameter, the accreting matter forms a thin disk at very small radii. We provide some analytical arguments to strengthen the numerical results; in particular, we estimate the radius where the gravitational force changes from attractive to repulsive and the critical value $|a|/M \\approx 1.4$ separating the two qualitatively different regimes of accretion. We briefly discuss the observational signatures which could be used to look for such exotic objects in the Galaxy and/or in the Universe. ",
        "introduction": "It is widely believed that the final product of the gravitational collapse of matter is a black hole (BH). In classical general relativity (GR), astrophysical BHs should be completely characterized by just three parameters: the mass $M$, the charge $Q$,  and the spin $J$. In this paper we focus on chargeless BHs.  The spin is often replaced by the Kerr parameter $a = J/M$. In classical GR, the values of $M$ and $a$ cannot be completely arbitrary, as they must satisfy the relation $|a| < M$, which is the condition for the existence of the horizon.  To see this we can examine the 3+1 dimensional Kerr solution. The position of the horizon is given by the expression~\\cite{mtw,lppt} \\be r_{H} = M + \\sqrt{M^2 - a^2} \\, . \\label{r-hor} \\ee It is clear that in (3+1)D spacetime the horizon cannot be formed if \\be M < |a| \\, . \\label{mass-limit} \\ee In the absence of a horizon, there would be naked singularities which are not allowed in GR. Indeed, if condition~(\\ref{mass-limit}) is fulfilled, the Kerr metric makes it possible to reach the physical singularity at $r=0$ from some large $r$ in finite time without crossing any horizon. One could thus consider closed time-like curves and violate causality (see e.g. section~66c of~\\cite{chandra} or ref.~\\cite{carter}). For this reason, usually some kind of cosmic censorship is assumed and naked singularities are forbidden~\\cite{penrose}. In particular, it is believed that naked singularities cannot be created by any physical process and therefore that the end-state of the gravitational collapse of matter is a Kerr BH with $|a| < M$~\\cite{penrose}.   However, in this paper we consider objects which {\\it do} violate the Kerr bound, i.e. with $|a|>M$. We call them ``super-spinars'', as proposed in~\\cite{horava}: since they have no event horizon, by the standard definition they are not BHs. Our main motivation is simple. The singularity can be viewed as the place where new physics should be expected: here observer-independent quantities like the scalar curvature diverge, while GR presumably breaks down above the Planck scale. It is therefore not unreasonable to expect that causality is conserved, not because the collapsing matter can form only objects with $|a| < M$, but because actually the central singularity is replaced by some high curvature region due to some quantum gravity effects, see e.g.~\\cite{hn04,horava}. In this case, there is apparently no reason to believe that the final product of the gravitational collapse of matter cannot have $|a| > M$.  Another possibility is that the collapsing matter forms a super-compact star with $|a| > M$ and exotic equations of state: now there is no central singularity, since the Kerr metric is a solution of Einstein equations only in vacuum; matter could have very exotic equation of state once it reaches densities so high that our knowledge of physics becomes inadequate. Actually, in general the metric at very small radii may deviate from the Kerr solution (the uniqueness theorem does not hold in absence of a regular horizon~\\cite{robinson}), but in our discussion we will neglect such a possibility.   In this paper, we extend the studies started in~\\cite{bf09, bft09}. The goal is to examine the main differences between the cases $|a| < M$ and $|a| > M$. Previous papers~\\cite{bf09, bft09} discussed implications on the apparent shape. There it was found that, even if the bound is violated by a small amount, the shadow cast by the super-spinar (i.e. how it blocks light coming to us from an object behind it) changes significantly from the case with $|a| < M$: the shadow for the super-spinar is about an order of magnitude smaller as well as distorted.  This distinction can be used as an observational signature in the search for these objects.  Based on recent observations at mm wavelength of the super-massive BH candidate at the Center of the Galaxy~\\cite{doeleman}, the authors speculated on the possibility that it might violate the Kerr bound.   In this paper we discuss the process of accretion in a Kerr background with arbitrary value of the Kerr parameter. For $|a|$ moderately larger than $M$, we find that the accreting matter cannot reach the central object, but is accumulated around it, forming a high density cloud. That may have interesting observational consequences. First, because of the high density and the high temperature of the plasma, the radiation produced in the accretion process can be much harder and more intense than the one coming from BHs. Second, there might be violent phenomena like bursts, when the amount of accumulated gas is large enough to gravitationally collapse. For higher values of $|a|$, the cloud evolves into a sort of disk, which is however very different from the usual disk of accretion around a BH: here the disk extends from $r \\approx M$ to the center, leading to rapid accretion, increasing efficiently the mass of the central object, and producing hard radiation at very small radii.  Unlike the case of Kerr BHs, we do not know if super-spinars are stable under small perturbations. Previous work has found that some very rapidly rotating objects in { 3+1} and higher dimensions can be unstable~\\cite{emparan, cardoso}. To address this point regarding the super-spinars studied in this paper, one should do a linearized analysis of the perturbations of these objects. However, the conclusion would be determined by the boundary conditions at the surface of super-spinars, which presumably depend on the quantum theory of gravity and are therefore unknown. Such a question cannot thus be addressed at present: here we just assume that super-spinars are stable and we study the accretion process onto these objects.    {\\it Conventions}: We use natural units $G_N = c = k_B = 1$. The metric has signature $(-+++)$. ",
        "conclusions": "In this paper we have studied the accretion process onto a spinning object. In particular, we have considered a Kerr spacetime with absolute value of the Kerr parameter $a$ (ratio of spin to mass) either smaller or larger than $M$, the mass of the object. In the first case, the spacetime contains a BH, in the second one a naked singularity. Our main motivation for considering the possibility of a naked singularity is based on the observation that the singularity is more likely a pathology of classical GR and that in the full theory it must be replaced by something else. We do not know how the central singularity is resolved, but our results are probably not significantly affected by the details of the correct theory. The only relevant quantity astrophysically is likely to be $|a|/M$.   We can distinguish three cases: \\begin{enumerate} \\item BHs with $|a|/M<1$. We find the usual accretion picture: injected matter always reaches a quasi-steady state configuration, in which matter is lost behind the event horizon at the same rate as it enters into the computational domain. \\item Super-spinars with $1<|a|/M<1.4$. Here the gas cannot reach the central object, because of a repulsive force in the neighborhood of the center. As a result, the gas is efficiently accumulated around the super-spinar. That leads to the formation and growth of a high density cloud. However, the accumulation process will stop at some point. One possibility is that it is interrupted by violent events due to the gravitational collapse of the cloud onto the object. This could be associated with the formation of a new object, either a BH or a heavier super-spinar. Another possibility is that the accumulated matter creates an event horizon, hiding the object from the outside, and with no abundant release of energy. \\item Super-spinars with $|a|/M \\gtrsim 1.4$. Now the repulsive force around the center is no longer capable of preventing a regular accretion of the object. Our simulations find that the flow forms a high density thin disk on the equatorial plane and reaches a quasi-steady state, i.e. matter enters and leaves the computational domain at the same rate. This disk is much closer to the center of the object than in the case of a standard BH. \\end{enumerate}"
    },
    "0910/0910.3574.txt": {
        "abstract": "We present a catalogue of variable stars in the near-infrared wavelength detected with overlapping regions of the 2MASS public images, and discuss their properties. The investigated region is in the direction of the Galactic center ($-30^\\circ \\lesssim l \\lesssim 20^\\circ, |b| \\lesssim 20^\\circ$), which covers the entire bulge. We have detected 136 variable stars, of which 6 are already-known and 118 are distributed in $|b| \\leq 5^\\circ$ region. Additionally, 84 variable stars have optical counterparts in DSS images. The three diagrams (colour-magnitude, light variance and colour-colour diagrams) indicate that most of the detected variable stars should be large-amplitude and long-period variables such as Mira variables or OH/IR stars. The number density distribution of the detected variable stars implies that they trace the bar structure of the Galactic bulge. ",
        "introduction": "The variable stars are useful tracers to study properties of the Galactic bulge. Among variables, the long-period variable stars (such as Semi-Regular Variables (SRVs), Miras and OH/IR stars) are detectable even in highly obscured region of the bulge owing to their high luminosities. These variable stars have been often used for investigating the Galactic bulge :e.g., measurement of the distance to the Galactic centre using Mira variables \\citep{Glass1995-MNRAS,Catchpole1999-IAUS,Groenewegen2005-AA,Matsunaga2009-pre}, a study of population of the bulge \\citep{Glass2001-MNRAS,Groenewegen2005-AA}. Therefore, a discovery of variable stars leads to provide us with important information about both variable star itself and the Galactic bulge.  Due to a large amount of dust in the bulge, past searches for variable stars have been mainly focused on relatively unobscured fields such as Baade's Windows \\citep[e.g., ][]{Evans1976-MNRAS,Glass1995-MNRAS}. Near-infrared observation is also better for collecting variable stars in the inner bulge, since near-infrared light suffers relatively small interstellar absorption. \\citet{Glass2001-MNRAS} detected large-amplitude variable stars in a $24 \\times 24$ arcmin$^2$ area of the Galactic centre from K-band survey, and obtained $\\sim 400$ objects with periods and amplitudes. \\citet{Schultheis2000-AA} looked for variable stars by comparing only twice DENIS observations, and presented two catalogues with the property of their variable star candidates. However, past surveys have only investigated small restricted area (up to a few deg$^2$).  Various imaging surveys have been performed in this two decades. In the imaging surveys, overlapping regions are produced for observing the sky without spatial gaps. In other words, there are regions observed multiple times in different epochs even if an imaging survey was performed only one time. Therefore, variable objects are expected to be detected with the overlapping regions. As claimed by \\citet{Schultheis2000-AA}, $\\sim 40\\%$ of variable stars can be recovered based on only twice-epoch magnitudes. Accordingly, the overlapping regions can be useful tools to search for variable stars. However, they have not been used in order to search for variable stars. When we can establish the search method based on overlapping region, it is also possible to search for variable stars using other survey data because survey images always contain overlapping regions. Furthermore, we can easily collect many variable stars in widely region (over a few tens of deg$^2$) with smaller time though we can extract less information compared with variable stars derived by an ordinary search (e.g., period, amplitude).  In this paper, we present a catalogue of variable stars obtained from overlapping region of 2MASS public images. In Sect. \\ref{Data}, we introduce 2MASS observation and data, and discuss the photometric accuracy in an overlapping region. In Sect. \\ref{Data-Analysis}, after introducing the detection procedure and criteria, we show a result of search and estimate the interstellar extinctions. In Sect. \\ref{Probability}, we discuss detection probability based on twice-epoch magnitudes. In Sect. \\ref{Discussion}, we discuss the near-infrared properties of the detected variable stars and the spatial distribution in the bulge.    %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "We have discovered variable stars using overlapping regions in the 2MASS public images. As a result of the investigation toward the Galactic centre ($-30^\\circ \\lesssim l \\lesssim 20^\\circ, |b| \\lesssim 20^\\circ$), we have detected 136 variable stars. Among which, 6 variable stars are already-known and 84 are accompanied by an optical counterpart in DSS images. The optical counterparts, however, are too faint to measure magnitudes (i.e., their brightness are nearby limiting magnitude of DSS images). The interstellar extinctions are estimated on the basis of the position of the upper giant branch, which are consistent with values of \\citet{Dutra2003-MNRAS}. Photometric properties of the detected variable stars indicate that most of them are large-amplitude AGB variables such as Mira variables and OH/IR stars in the Galactic bulge. Additionally, the bar-like structure of the bulge is detected by the number density distribution of the detected variable stars.  This paper demonstrates that the search for variable stars using overlapping regions is an useful method. In this paper, there is a small number of detection due to the strict detection criteria, but it is possible to detect a lot of variable stars if we can judge the variability more precisely with tender criteria. Establishing this search method, we can discover variable stars in widely region. If extracting more variable stars in the bulge, the structure of the bulge might be revealed using them as tracers, since the detected variable stars in this paper trace the structure of the bulge."
    },
    "0910/0910.1544_arXiv.txt": {
        "abstract": "{We present observations of two new single-lined eclipsing binaries, both consisting of an Am star and an M-dwarf, discovered by the Wide Angle Search for Planets transit photometry survey.  Using WASP photometry and spectroscopic measurements we find that HD186753B has an orbital period of $P=1.9194$ days, a mass of $M=0.24\\pm0.02 M_{\\odot}$ and radius of $R=0.31^{+0.06}_{-0.06} R_{\\odot}$; and that TCY7096-222-1B has an orbital period of $P=8.9582$ days, a mass of between 0.29 and 0.54 $M_{\\odot}$ depending on eccentricity and radius of $R=0.263^{+0.02}_{-0.07} R_{\\odot}$.  We find that the Am stars have relatively low rotational velocities that closely match the orbital velocities of the M-dwarfs, suggesting that they have been ``spun-down'' by the M-dwarfs.} ",
        "introduction": "The radius of a star is one of its most fundamental properties, yet for sub-solar masses models have not been able to provide accurate radius predictions.  Citing the discrepancies between model and empirical radius measurements, Chabrier et al. (\\cite{chabrier}) found that large surface spot coverage decreases the photospheric temperature.  The star compensates by increasing its radius to conserve radiative pressure.  This was confirmed by an empirical activity-radius study by Lopez-Morales (\\cite{lopez-morales}).  In addition, Berger et al. (\\cite{berger}) found a correlation between an increase in metalicity and a larger-than-expected radius.  Because of their low intrinsic brightness, low-mass stars (LMS) are particularly difficult to study. LMS in eclipsing binary systems (EBLM), however, provide a direct way to obtain radius measurements and are therefore a valuable tool for testing models of stellar structure in the low-mass region.  A by-product of wide-angle transit photometry planet-searching projects is the discovery of EBLMs (e.g. Fernandez et al. \\cite{newtrespaper}).  The metallic-line Am stars are a class of peculiar A-type stars that are slow rotating, thought to have had their rotational velocity reduced by a near stellar companion.  Spectroscopic orbits of many Am stars have been reported (e.g. Carquillat \\& Prieur \\cite{ampaper}, Renson \\& Manfroid \\cite{renson}) and LMS are thought to be responsible for reducing the rotational velocity of Am stars (Carquillat \\& Prieur \\cite{ampaper}). Here we report the discovery of two single lined A-M binaries, HD186753 and TYC7096-222-1, the first EBLMs discovered from the Wide Angle Search for Planets (WASP) planet-hunting project.  There are four likely eclipsing A-M systems that have previously been announced; three of these systems were announced in Dreizler et al. (2002) and one in Pont et al. (2005). ",
        "conclusions": "\\label{discussion} The rotational velocity of A8 stars is expected to be $\\sim~200\\rm{km\\ s^{-1}}$ (Gray \\cite{gray}).  The relatively low rotational velocities of HD186753A and TYC7096-222-1A, $v\\sin{i}=65.0\\pm5.0$ and $35.0\\pm5.0 \\rm{km\\ s^{-1}}$, respectively, suggests that they have been ``spun-down'' by their M-dwarf companions.  The stellar rotational angular velocity of HD186753A is $(3.67\\pm0.65)\\times10^{-4} \\rm{rad\\ s^{-1}}$ which is greater than the M-dwarf orbital angular velocity of $(3.79\\pm0.01)\\times10^{-5} \\rm{rad\\ s^{-1}}$. The rotational angular velocity of TYC7096-222-1A is $(3.16\\pm0.47)\\times10^{-5} \\rm{rad\\ s^{-1}}$ which is also greater than the M-dwarf orbital angular velocity of $(8.12\\pm0.02)\\times10^{-6} \\rm{rad\\ s^{-1}}$.  The synchronisation timescales are 2.8 Myr for HD186753 and 0.92 Gyr for TYC7096-222-1 (Zahn 1977), suggesting that each system is younger than the synchronisation time.  Values of $\\log{g}$ can be used as an age estimate, but our value is not reliable enough.  Higher SNR spectra is required to determine the age of the systems.  The eccentricity of HD186753B, $e=0.269\\pm0.087$, is quite high for a short period binary and could also indicate a young stellar age as the circulization time is 95 Myr (Zahn 1977).  The circulization time for  TYC7096-222-1B is 295 Gyr.  If TYC7096-222-1B has a mass of $M_{2}=0.54\\pm0.06 M_{\\odot}$ (using the $e=0.75$ model) then we would expect a secondary eclipse of $0.003\\pm1\\times{}10^{-4}$ mag deep.  There appears to be no sign of this in the WASP photometry at any phase.  TYC7096-222-1B therefore has a mass between $0.29-0.54 M_{\\odot}$ depending on the eccentricity. The certainty of $e$ for both objects could be improved with greater radial velocity phase coverage.  The value of $e$ for HD186753B could be explained by a tertiary component either in the system or by a recent near-miss, although we find no evidence in our data for either scenario.  Both primaries show supersolar Fe abundances, an underabundance of Ca and Sc and overabundances of Y and Ba. These abundances are typical of Am stars (e.g. Wolff \\cite{wolff}; Hundt \\cite{hundt}).  Carquillat \\& Prieur (\\cite{ampaper}) found that the mean companion mass to an Am spectral type is $0.8\\pm{0.5}M_{\\odot}$ with a mean orbital period of 1.995 days. By these stellar properties HD186753 and TYC7096-222-1 are fairly typical Am systems.  The radii of stars with masses below $0.3-0.35\\rm{M_{\\odot}}$ agree well with the Baraffe et al. (\\cite{baraffe}) isochrone models (Lopez-Morales \\cite{lopez-morales}).  The mass-radius relation TYC7096-222-1B agree within $1\\sigma$ of the Baraffe et al. (\\cite{baraffe}) isochrones, as shown in Fig. \\ref{m-r}. Whilst the mass-radius relation of HD186753B, however, is only just within $1\\sigma$ of the Baraffe et al. (\\cite{baraffe}) isochrones.  M-dwarfs in binaries are found to be more active than solitary M-dwarfs (Lopez-Morales \\cite{lopez-morales}). Increased activity causes a decrease in photospheric temperature, which then causes an increase in radius to conserve radiative pressure (Ribas \\cite{Ribas}).  X-ray activity is an indicator of stellar activity, although no X-ray data has been published on either object.  Being in a tight orbit, the photosphere of HD186753B is strongly irradiated by HD186753A.  The radiation intensity received from HD186753A at the photosphere of HD186753B, and also therefore the radiative pressure, is $1.32\\pm0.02$ times the radiation generated; the value for TYC7096-222-1 is lower at $0.72\\pm0.04$.  Higher precision radial velocity and photometry than the values given here are required to ascertain as to whether this extra source of heating, or increased levels of activity are affecting the mass-radius relation of the M-dwarfs.  \\begin{figure} \\centering \\resizebox{\\columnwidth}{!}{\\includegraphics[angle=270,bb=50 50 554 770]{M_vs_R.ps}} \\caption{Mass-radius relationship for solitary stars, EBLMs, HD186753B and TCY7096-222-1 superimposed on 5 Gyr, [M/H]=0, $L_{\\rm{mix}}=1H_{p}$ Baraffe et al. (\\cite{baraffe}) isochrones. The thick black line indicates the range of masses for TYC7096-222-1B depending on the eccentricity. M-dwarfs with uncertainties in mass $>0.03 M_{\\odot}$ have been omitted from the figure for illustrative clarity.} \\label{m-r} \\end{figure}   \\begin{table*} \\begin{tabular}{llll} Parameter & HD186753B & TYC7096-222-1B ($e=0.00$) & TYC7096-222-1B ($e=0.75$) \\\\   \\hline\\hline Eclipse epoch (HJD) & $3940.40144^{+0.00091}_{-0.00081}$ & $4373.01637^{+0.00066}_{-0.00087}$ & - \\\\ Orbital period (days) & $1.9193851^{+0.0000412}_{-0.0000393}$ & $8.9582591^{+0.0000346}_{-0.0000314}$ &  - \\\\ Eclipse duration (days) & $0.1662^{+0.0038}_{-0.0031}$ & $0.2430^{+0.0039}_{-0.0042}$  &  - \\\\ Secondary/primary area ratio, $(R_{2}/R_{1})^2$ & $0.0148^{+0.0005}_{-0.0003}$ & $0.0251^{+0.0006}_{-0.0005}$ &  - \\\\ Impact parameter, $b (R_{*})$ & $0.264^{+0.191}_{-0.154}$ & $0.250^{+0.157}_{-0.142}$  &  - \\\\ Stellar reflex vel., $\\rm{K_{1}}$ ($\\rm{km\\ s^{-1}}$) & $-27.449\\pm1.751$ & $-20.341\\pm2.974$  & $-58.432\\pm6.122$ \\\\ Centre-of-mass vel., $\\gamma$ ($\\rm{km\\ s^{-1}}$) & $-14.641\\pm2.980$ & $4.074\\pm1.533$  & $8.776\\pm0.748$ \\\\ Orbital separation, $a$ (AU) & $0.0370^{+0.0012}_{-0.0013}$ & $0.0990^{+0.0031}_{-0.0034}$  &  - \\\\ Orbital inclination, $i$ (deg) & $87.09^{+1.69}_{-1.79}$ & $89.00^{+0.57}_{-0.73}$  &  - \\\\ Orbital eccentricity, $e$ & $0.269\\pm0.087$ & (fixed=0.00)  & (fixed=0.75) \\\\ Longitude of periastron, $\\omega$ (deg) & $166.7\\pm5.9$ & (not fitted)  & $167.2\\pm61$ \\\\ Stellar Mass, $M_{2}(M_{\\odot})$ & $0.236\\pm0.017$ & $0.286\\pm0.019$  & $0.544\\pm0.057$ \\\\ Stellar Radius, $R_{2}(R_{\\odot})$ & $0.307^{+0.057}_{-0.057}$ & $0.263^{+0.020}_{-0.071}$  &  - \\\\ Luminosity ratio, $L_{1}/L_{2}$ & $979.7$ & $520.4$  & $361.9$ \\\\ \\end{tabular} \\caption{\\label{paramtable} Parameters of HD186753B and TYC7096-222-1B and their orbits.  The parameters of TYC7096-222-1B determined from the photometry are the same for both eccentricities.} \\end{table*}"
    },
    "0910/0910.4966_arXiv.txt": {
        "abstract": "We utilize color information for an HI-selected sample of 195 galaxies to explore the star formation histories and physical conditions that produce the observed colors.  We show that the HI selection creates a significant offset towards bluer colors that can be explained by enhanced recent bursts of star formation. There is also no obvious color bimodality, because the HI selection restricts the sample to bluer, actively star forming systems, diminishing the importance of the red sequence. Rising star formation rates are still required to explain the colors of galaxies bluer than $g-r <$ 0.3.  We also demonstrate that the colors of the bluest galaxies in our sample are dominated by emission lines and that stellar population synthesis models alone (without emission lines) are not adequate for reproducing many of the galaxy colors.  These emission lines produce large changes in the $r-i$ colors but leave the $g-r$ color largely unchanged. In addition, we find an increase in the dispersion of galaxy colors at low masses that may be the result of a change in the star formation process in low-mass galaxies. ",
        "introduction": "The star formation history, metallicity, and current star formation rate all contribute to the observed colors of a galaxy.  Remarkably, these factors work in concert to yield a fairly well defined locus in galaxy color-color space (Strateva et al.~2001).  Deviations from this locus as well as the morphology of the locus itself, can lead to significant insight into the underlying processes taking place within galaxies.  Many previous studies have utilized the broad-band colors of galaxies to investigate the star formation histories and metallicities of galaxies (e.g. Tinsley 1972; Searle et al.~1973; Tinsley \\& Gunn 1976; Balcells \\& Peletier 1994; Roberts \\& Haynes 1994; de Jong 1996; Bell \\& de Jong 2000; Galaz et al.~2002; Gavazzi et al.~2002; Bell et al.~2003; MacArthur et al.~2004; Zackrisson, Bergvall \\& Ostlin 2005; Driver et al.~2006; Skibba et al.~2008).  The advent of large surveys such as the Sloan Digital Sky Survey (SDSS; York et al.~2000) and the Two Micron All Sky Survey (2MASS; Skrutskie et al.~2006) has created a wealth of uniform data and the statistical foothold to investigate the bimodality of galaxies (Baldry et al.~2004; Kauffmann et al.~2003, 2004), luminosity function (Blanton et al.~2001; Ball et al.~2006), and the average properties of nearby galaxies (Blanton et al.~2003a, 2003b, 2005; Geha et al.~2006; Maller et al.~2008; Skibba et al.~2008). While the large optical and infrared surveys have made large contributions to our understanding of galaxy evolution, they only trace the stellar component of galaxies and do not trace other baryonic material such as cold gas.  Galaxies in the local universe span a range of star formation histories -- from blue, gas-rich, low-surface-brightness (LSB) galaxies that are slowly turning their gas into stars with low star forming efficiencies, to red, gas-poor galaxies, that have formed the bulk of their stars in the past.  Stars dominate the visible light output of most galaxies, and thus galaxies detected by traditional optical or infrared imaging have well developed stellar populations.  In contrast, the natural way to identify gas-rich, less evolved galaxies is by their 21 cm HI radio emission.  Aside from its importance for global star formation, a sample of galaxies with both gaseous and stellar information allows for a more complete census of the local baryons (a constraint vital to the calibration of n-body simulations; Governato et al.~2007; Brooks et al.~2009).  Previous studies have combined large HI surveys with optical and infrared samples, namely the Arecibo Duel Beam and Slice Surveys with the Two Micron All Sky Survey (2MASS; Jarrett et al.~2000; Rosenberg et al.~2005) and the merging of HIPASS with SuperCOSMOS (Hambly et al.~2001; Hambly, Irwin \\& MacGillivary 2001b; Hambly et al.~2001c; Doyle et al.~2005).  Rosenberg et al.~(2005) were able to probe the baryonic content of a large sample of galaxies, but were limited by the shallow depth of 2MASS, which does not have data for many of the LSB galaxies in the sample.  The HIPASS/SuperCOSMOS sample of Doyle et al.~(2005) contains optical data for more than 3600 HI selected galaxies but also suffers from the shallow depth of the SuperCOSMOS optical data.  Recent studies have combined the Parkes HI Equatorial Survey (ES) with the SDSS (Disney et al.~2008; Garcia-Appadoo et al.~2009; West et al.~2009; hereafter W09).  The deep SDSS optical data provides information about the stellar content for \\emph{all} ES HI sources where the two surveys overlap.  In addition, the uniform, accurate, and well-calibrated photometry of the HI-selected ES/SDSS sample allows for a more detailed exploration of the factors affecting galaxy colors, particularly for galaxies with large reservoirs of gas.  The optical colors are reasonably sensitive to age, although IR colors are needed to constrain metallicity (Bell et al.~2003).  There is thus an unavoidable degeneracy between age and metallicity when using only the optical colors available with SDSS (Bell \\& de Jong 2001; Bell et al.~2003).  Our sample does not have a complete set of near IR counterpart data and some of our results will reflect this limitation.  In this paper, we briefly describe the ES/SDSS sample in \\S2 and examine the colors of HI-selected galaxies by comparing them to stellar population synthesis models (\\S3.1) and the colors of optically selected galaxies (\\S3.2).  We also investigate how line emission affects the broadband colors of gas-rich galaxies (\\S3.3). We demonstrate an increased dispersion in the colors of galaxies at low masses, and investigate its possible origins (\\S3.4).  We summarize and discuss our results in \\S4. ",
        "conclusions": "We used the population synthesis models of Bruzual \\& Charlot (2003) to model the SFHs and metallicities of galaxies in the ES/SDSS sample. We found that red galaxies have super-solar metallicities and have SFHs that are consistent with them forming the bulk of their stars in the distant past.  These red galaxies have presumably exhausted their early gas supply but have recently acquired gas through mergers and infall. Although the infalling gas is likely low metallicity and could in principle dilute the metallicity, the small amount of gas in red galaxies (low gas fractions) is not enough to make an observable difference in the average metallicity of the ES/SDSS systems.  Bluer galaxies have lower metallicities and their mean stellar ages are younger (increasing $\\tau$ values).  The gas fractions suggests that these blue systems have at least as much gas as they do stars. Previous studies have indicated that these types of galaxies are less massive and have disks that are stable against collapse, making their star formation much less efficient.  They have also likely retained much of their initial HI and are slowly converting it into stars as is suggested by the large (or negative) $\\tau$ values.  As galaxies pass the threshold for disk stability, their star formation becomes sporadic and there is a clear dispersion in the colors due to variations in burst age, burst strength and superimposed emission lines.  The idealistic Bruzual and Charlot models are likely only loose guides for these systems, as their star formation histories no longer follow smooth exponential functions.  Figure \\ref{bc1} indicates that there are a few galaxies in the regime of exponentially increasing SFR (negative $\\tau$ values).  This is not surprising as there are few mechanisms that will increase the SFR of a galaxy.  The likely culprit for these systems is the infall of gas.  As gas is accreted, the gas densities in the galaxies will increase and local gravitational collapse will become more efficient.  This will increase the SFR as well as the eventual metallicity of the galaxy, pushing its $g-r$ colors to the blue (more recent star formation) and its $r-i$ colors slightly to the red (more metals).  We also showed that HI selected galaxies are offset from the SDSS galaxy locus, especially at the red end, and that this is likely due to bursts of star formation in the past few hundred Myr.  The gas that induces these recent bursts is not primordial and is best explained by the accretion of gas rich dwarfs.  The bluest galaxies in the ES/SDSS sample are not explained by population synthesis models alone.  Their colors can be modeled only with the inclusion of emission lines.  We showed that emission line spectra with reasonable SFRs can explain the colors of the bluest galaxies in our sample.  We also showed that the distribution of galaxies at the red end of the color-color locus has a very small dispersion that continues to rise into the blue regime. This change in dispersion appears to correlate with stellar mass, rotation velocity, surface brightness and especially stellar surface density. It is possible that the change in dispersion is evidence for a significant change in the way stars form in galaxies at a given mass scale.  Massive galaxies have unstable disks and efficiently convert most of their primordial reservoirs of gas into stars in the first few Gyrs after their formation.  These galaxies are bulge dominated and have redder, older and more massive populations of stars.  We do see a sharp transition in the color distribution at a rotation velocity of 80 km~s$^{-1}$.  This transition may be related to previous results that have found sharp transitions in galaxy properties as a function of rotational velocity (Dalcanton et al.~2004; Lee et al.~2007).  However, further investigations are required to explore the various velocity transitions as a function of galaxy properties.  It may also be the case that the increase in dispersion is nothing more than a selection effect related to the surface brightness. As the surface brightness of a system decreases, the effect of a single burst of stars on the color becomes increasingly large, possibly explaining the increase in scatter.  This effect might explain the onset of emission line dominated colors in the bluest galaxies as they can influence of the colors of galaxies devoid of massive stellar populations.  We note that the change in color dispersion is not easily seen in the volume selected SDSS data.  Applying the HI selection identifies a low dispersion subsample of the SDSS galaxies, most notably at the red end.  As mentioned above, the red HI selected galaxies appear to be bluer in $g-r$ than the ``main'' SDSS sample. These two features are likely related.  It is possible that the large distribution in color at the red end of the SDSS main sample is due to the diversity of gas content.  Because most of the red galaxies in SDSS exhausted their original supply of gas long ago, this dispersion may be an indication of the spread in time since the last major gas infall.  We leave further discussion of the ``color offset'' to future study."
    },
    "0910/0910.3893_arXiv.txt": {
        "abstract": "Sagittarius~A* is the source of near infrared, X-ray, radio, and (sub)millimeter emission associated with the supermassive black hole at the Galactic Center.  In the submillimeter regime, Sgr~A* exhibits time-variable linear polarization on timescales corresponding to $< 10$ Schwarzschild radii of the presumed $4 \\times 10^6$~M$_\\sun$ black hole.  In previous work, we demonstrated the potential for total-intensity (sub)millimeter-wavelength very long baseline interferometry (VLBI) to detect time-variable -- and periodic -- source structure changes in the Sgr~A* black hole system using nonimaging analyses.  Here we extend this work to include full polarimetric VLBI observations.  We simulate full-polarization (sub)millimeter VLBI data of Sgr~A* using a hot-spot model that is embedded within an accretion disk, with emphasis on nonimaging polarimetric data products that are robust against calibration errors. Although the source-integrated linear polarization fraction in the models is typically only a few percent, the linear polarization fraction on small angular scales can be much higher, enabling the detection of changes in the polarimetric structure of Sgr~A* on a wide variety of baselines.  The shortest baselines track the source-integrated linear polarization fraction, while longer baselines are sensitive to polarization substructures that are beam-diluted by connected-element interferometry.  The detection of periodic variability in source polarization should not be significantly affected even if instrumental polarization terms cannot be calibrated out.  As more antennas are included in the (sub)mm-VLBI array, observations with full polarization will provide important new diagnostics to help disentangle intrinsic source polarization from Faraday rotation effects in the accretion and outflow region close to the black hole event horizon. ",
        "introduction": "The Galactic Center source Sagittarius~A* (Sgr~A*) provides the best case for high-resolution, detailed observations of the accretion and outflow region surrounding the event horizon of a black hole.  There are several compelling reasons to observe Sgr~A* with very long baseline interferometry (VLBI) at (sub)millimeter\\footnote{We shall henceforth use the term ``millimeter'' to denote wavelengths of 1.3~mm of shorter (in contrast with observations at 3~mm and 7~mm, which are sometimes also referred to as ``millimeter'' wavelengths).} wavelengths.  The spectrum of Sgr~A* peaks in the millimeter \\citep[][and references therein]{markoff07}.  Interstellar scattering, which varies as the wavelength $\\lambda^2$, becomes less than the fringe spacing of the longest baseline available to VLBI in the millimeter-wavelength regime.  Indeed, VLBI on the longest baselines available at 345~GHz probes scales of twice the Schwarzschild radius ($R_\\mathrm{S}$) for a $4 \\times 10^6$~M$_\\sun$ black hole.  From previous observations at 230~GHz, it is known that there are structures on scales smaller than a few $R_\\mathrm{S}$ \\citep{doeleman08}.  Such high angular resolution, presently unattainable by any other method (including facility instruments such as the Very Long Baseline Array), is necessary to match the expected spatial scales of the emitting plasma in the innermost regions surrounding the black hole and will be required to unambiguously determine the inflow/outflow morphology and permit tests of general relativity.  This sensitivity to small spatial scales also makes millimeter polarimetric VLBI possible.  Although the linear polarization fraction of Sgr~A* integrated over the entire source is only a few percent \\citep[e.g.,][]{marrone07}, the fractional polarization on small angular scales is likely much larger.  In general, relativistic accretion flow models predict that the electric vector polarization angle (EVPA) will vary along the circumference of the accretion disk \\citep{bromley01,broderick05,broderick06}, indicating that single-dish observations and connected-element interferometers probably underestimate linear polarization fractions due to beam depolarization.  Polarized synchrotron radiation coming from Sgr~A* was detected by \\citet{aitken00} at millimeter and submillimeter wavelengths. Multiple observations since then have demonstrated that the polarized emission is variable on timescales from hours to many days \\citep{bower05,macquart06,marrone06a,marrone07,marrone08}.  In one case, the timescale of variability and the trace of polarization in the Stokes $(Q,U)$ plane of the millimeter-wavelength emission are suggestive of the detection of an orbit of a polarized blob of material \\citep{marrone06b}.  Near infrared observations by \\citet{trippe07} are also consistent with a hot spot origin for periodic variability.  It is possible that connected-element interferometry may suffice to demonstrate polarization periodicity, but millimeter-wavelength VLBI, which effectively acts as a spatial filter on scales of a few to a few hundred $R_\\mathrm{S}$, can be more sensitive to changing polarization structures.  Initial millimeter VLBI observations of Sgr~A* will necessarily utilize non-imaging analysis techniques, for reasons outlined in \\citet[][henceforth Paper~I]{doeleman09}.  One way to do this is to analyze so-called interferometric ``closure quantities,'' which are relatively immune to calibration errors \\citep{rogers74,rogers95}.  In Paper~I, we considered prospects for detecting the periodicity signature of a hot spot orbiting the black hole in Sgr~A* via closure quantities in total-intensity millimeter-wavelength VLBI.  In the single polarization case, it is necessary to construct closure quantities from at least three or four antennas in order to produce robust observables, since the timescales of atmospheric coherence and frequency standard stability do not permit standard nodding calibration techniques.  Closure quantities can be used in full-polarization observations as well, but it is also possible to construct robust observables on a baseline of two antennas by taking visibility ratios between different correlation products.  In this work, we extend our techniques to explore polarimetric signatures of a variable source structure in Sgr~A*, with emphasis on ratios of baseline visibilities. ",
        "conclusions": "(Sub)millimeter-wavelength VLBI polarimetry is a very valuable diagnostic of emission processes and dynamics near the event horizon of Sgr~A*.  We summarize the findings in this paper as follows:  \\begin{itemize} \\addtolength{\\itemsep}{-0.8ex} \\item Millimeter-wavelength polarimetric VLBI can detect changing source structures.  Despite low polarization fractions seen with connected-element interferometry, the much higher angular resolution data provided by VLBI will be far less affected by beam depolarization and contamination from dust polarization. Polarimetric VLBI provides an orthogonal way to detect periodic structural changes as compared with total intensity VLBI.  \\item Ratios of cross- to parallel-hand visibilities are robust baseline-based observables.  Short VLBI baselines approximately trace the integrated polarization fraction and position angle of the inner accretion flow of Sgr~A*, while longer VLBI baselines resolve smaller structures.  \\item Calibration of instrumental polarization terms is not necessary to detect a changing source structure, including periodicity, in Sgr~A*.  \\item Polarimetric VLBI may be able to disentangle the effects of rotation measure from intrinsic source polarization.  Initial results will likely come from observations of the timescale of polarimetric variability.  If the initial array is expanded to allow high-fidelity imaging, polarimetric VLBI may be able to map the Faraday rotation region and directly infer the density and magnetic field structure of the emitting region in Sgr~A*.   \\end{itemize}"
    },
    "0910/0910.5292_arXiv.txt": {
        "abstract": "We obtained SDSS spectra for a set of 37 radio-quiet quasars (RQQSOs) that had been previously examined for rapid small scale optical variations, or microvariability.  Their \\hbeta and \\mgii emission lines were carefully fit to determine line widths (FWHM) as well as equivalent widths (EW) due to the broad emission line components.  The line widths were used to estimate black hole masses and Eddington ratios, $\\ell$.  Both EW and FWHM are anticorrelated with $\\ell$.  The EW distributions provide no evidence for the hypothesis that a weak jet component in the RQQSOs is responsible for their microvariability. ",
        "introduction": "Over the past 15 years there have been rather extensive examinations of a significant sample of radio-quiet QSOs (RQQSOs) and Seyfert galaxies for small brightness changes (typically 0.02 mag) over short times (a few hours) (e.g., Stalin et al.\\ 2004b; Carini et al.\\ 2007; Ram{\\'i}rez et al.\\ 2009). This phenomenon of microvariability, or intranight optical variability (INOV) was first confirmed for blazars (e.g., Miller, Carini \\& Goodrich 1989; Carini 1990) for which microvariability almost certainly arises from the relativistic jet (e.g., Marsher, Gear \\& Travis 1992), at least when the source is in an active state; however, in low states it is possible that rapid fluctuations are due to processes originating in or just above the accretion disc (for a review, see Wiita 2006).  Since RQQSOs lack significant jets, the microvariability in radio-quiet objects may arise from processes on the accretion disc itself, and thus could possibly be used to probe the discs (e.g., Gopal-Krishna, Wiita \\& Altieri 1993; Stalin et al.\\ 2004a; Kelly, Bechtold \\& Siemiginowska 2009).  Recently Carini et al.\\ (2007) reported  new observations for several sources and also  compiled a sample of 117 objects from the literature which have been searched for microvariability (their Table 3). Of these, 47 are classified as Seyfert galaxies, 6 as broad absorption line (BAL) QSOs, and 64 as QSOs. In addition, in order to learn which classes might be the most fruitful for making further searchs for microvariability they also classified their sample in terms of following: redshift distribution; radio loudness, through $R$, the ratio of the radio [5 GHz] flux to the optical [4400\\AA] flux); optical magnitude, $m_V$; luminosity; and observing strategy. In their entire sample 21.4\\% of the objects were found to exhibit microvariability, but among objects classified as Seyfert galaxies, BAL QSOs and QSOs, microvariability was found in 17\\%, 50\\% and 23.2\\%, respectively (Carini et al.\\ 2007).  The observed high fraction of microvariations in BAL QSOs (although the sample is quite small) suggests that while planning microvariability studies to investigate the physical processes in or near the accretion disc, it might be worthwhile to invest more observing time on the BALQSO class. \\par  With improvements in observation quality as well as in sample size, constraints on the models capable of producing the intranight variability have also improved over last decade.  Recently, Czerny et al.\\ (2008) have used non-simultaneous optical and X-ray data of  10 RQQSOs with confirmed INOV to compare observational constraints on the variability properties with the predictions of theoretical models such as: (i) irradiation of an accretion disc by a variable X-ray flux (e.g., Rokaki, Collin-Souffrin \\& Magnan 1993; Gaskell 2006); (ii) an accretion disc instability (e.g., Mangalam \\& Wiita 1993); (iii) the presence of a weak Doppler boosted jet, or ``blazar component'' (e.g., Gopal-Krishna et al.\\ 2003). Their investigation suggests that a blazar component model yields the highest probability of detecting INOV.  In this blazar component scenario, spectral properties of the sources can play a crucial role in constraining the models further. For instance, if blazar components are dominating the variability of RQQSOs, then, due to the increase in the continuum level, one would expect emission lines to be diluted.  Therefore smaller equivalent widths (EWs) of prominent emission lines such as \\hbeta and \\mgii should be detected in sources that showed microvariability when compared to their average values in the whole sample. If we take the extreme case of BL Lacertae objects, which often lack observable emission lines and are usually defined as objects that have no emission line with an EW $\\geq$5\\AA \\ (e.g., Stickel et al.\\ 1991; cf., March{\\~a} et al.\\ 1996), this dilution by the jet component is understood to be severe.  So it becomes very important to test whether microvariability of RQQSOs has any correlation with spectral parameters such as EW and full width at half maximum (FWHM) of prominent emission lines. The \\hbeta and \\mgii emission lines are very promising for such investigations, as these lines have also been found to be very useful in estimating other key parameters of AGN central engines such as black hole (BH) mass (e.g., McGill et al.\\ 2008) and Eddington ratio (e.g., Dong et al.\\ 2009a, 2009b). The average \\hbeta EW of a large sample of quasars is found to be around 62.4\\AA \\ (Foster et al.\\ 2001), so any correlation of \\hbeta EW or FWHM {\\bf would} not only give insight about the nature of variability, such as the presence or absence of blazar components, but also {\\bf would} be very useful for making a promising sample for future microvariability studies.  Similarly the measurements of BH masses for RQQSOs that show microvariability may well be another important constraint {\\bf on} the models trying to understand the nature of their optical microvariability.  Here we have worked toward these goals by exploiting the optical spectra available from Sloan Digital Sky Survey (SDSS) Data Release 7 (DR7; Abazajian et al.\\ 2009) with careful spectral modeling of the \\hbeta and \\mgii emission line regions. First we aim to investigate any effect of these key spectral parameters (e.g., EW and FWHM) on microvariability of RQQSOs. Second we estimate other relevant AGN parameters such as the black hole mass, M$_{bh}$, and the Eddington ratio in the context of our RQQSOs sample, which has been extensively searched for microvariability. \\par  The paper is organized as follows. Section 2 describes the data sample and selection criteria while Section 3 gives details of our spectral fitting procedure. In Section 4 we focus on BH mass measurements and in Section 5 we give estimates of Eddington ratios and of BH growth times.  Section 6 gives a discussion and conclusions.  Throughout, we have used a cosmology with $H_{\\rm 0}$=70 km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_{\\rm M}$=0.3 and $\\Omega_{\\rm \\Lambda}$=0.7.  \\begin{table*} \\centering \\begin{minipage}{140mm} \\caption{ Our sample of radio-quiet QSOs and Seyfert galaxies from the compilation of Carini et al.\\ (2007)} \\label{lab:tabsamp} \\begin{tabular}{@{}llllllllllcl@{}} \\hline \\multicolumn{1}{c}{QSO Name\\footnote{Name from V{\\'e}ron-Cetty \\& V{\\'e}ron (2006)}}   & {$z_{em}$} & \\multicolumn{3}{c}{$\\alpha_{2000}$} & \\multicolumn{3}{c}{$\\delta_{2000}$} &{m$_{V}$} & {M$_{V}$} &  {Variable?} & {Class} \\\\ \\hline \\\\  Mrk 1014            & 0.163        &   01& 59& 50.2 &   $+$00& 23& 41&   15.9 &   $-$23.8 &   Y&  Sy 1    \\\\ US 3150              & 0.467        &   02& 46& 51.9 &   $-$00& 59& 31&   17.1 &   $-$25.0 &  N&  NLS1  \\\\ US 3472              & 0.532        &   02& 59& 37.5 &   $+$00& 37& 36&   16.8 &   $-$25.7 &  N&  QSO     \\\\ Q  J0751$+$2919      & 0.912        &   07& 51& 12.3 &   $+$29& 19& 38&   16.2 &   $-$27.7 &  Y&  QSO     \\\\ PG 0832$+$251        & 0.331        &   08& 35& 35.9 &   $+$24& 59& 41&   16.1 &   $-$25.5 &  Y&  QSO     \\\\ US 1420              & 1.473        &   08& 39& 35.1 &   $+$44& 08& 11&   17.5 &   $-$27.3 &  N&  QSO     \\\\ US 1443              & 1.564        &   08& 40& 30.0 &   $+$46& 51& 13&   17.2 &   $-$27.8 &  N&  QSO     \\\\ US 1498              & 1.406        &   08& 42& 15.2 &   $+$45& 25& 44&   17.7 &   $-$26.7 &  N&  QSO     \\\\ US 1867              & 0.513        &   08& 53& 34.2 &   $+$43& 49& 01&   16.9 &   $-$25.2 &  N&  Sy 1    \\\\ PG 0923$+$201        & 0.192        &   09& 25& 54.7 &   $+$19& 54& 04&   15.5 &   $-$24.8 &  Y&  Sy 1    \\\\ PG 0931$+$437        & 0.456        &   09& 35& 02.6 &   $+$43& 31& 11&   16.0 &   $-$26.4 &  N&  QSO     \\\\ PG 0935$+$416        & 1.966        &   09& 38& 57.0 &   $+$41& 28& 21&   16.8 &   $-$28.8 &  N&  QSO     \\\\ CSO 233              & 2.030        &   09& 39& 35.1 &   $+$36& 40& 01&   18.4 &   $-$27.3 &  N&  QSO     \\\\ CSO 18               & 1.300        &   09& 46& 36.9 &   $+$32& 39& 51&   17.0 &   $-$27.9 &  N&  QSO     \\\\ US 995               & 0.226        &   09& 48& 59.4 &   $+$43& 35& 18&   16.9 &   $-$23.4 &  Y&  QSO     \\\\ PG 0946$+$301        & 1.220        &   09& 49& 41.1 &   $+$29& 55& 19&   16.2 &   $-$28.2 &  N&  BAL QSO \\\\ CSO 21               & 1.190        &   09& 50& 45.7 &   $+$30& 25& 19&   17.3 &   $-$27.3 &  N&  QSO     \\\\ TON 34               & 1.925        &   10& 19& 56.6 &   $+$27& 44& 02&   15.7 &   $-$29.8 &  Y&  QSO     \\\\ PG 1049$-$005        & 0.357        &   10& 51& 51.5 &   $-$00& 51& 17&   15.8 &   $-$25.7 &  N&  Sy 1    \\\\ TON 52               & 0.434        &   11& 04& 07.0 &   $+$31& 41& 11&   17.3 &   $-$24.9 &  Y&  QSO     \\\\ PG 1206$+$459        & 1.155        &   12& 08& 58.0 &   $+$45& 40& 36&   15.7 &   $-$28.4 &  N&  QSO     \\\\ UM 497               & 2.022        &   12& 25& 18.4 &   $+$02& 06& 57&   17.7 &   $-$28.0 &  N&  QSO     \\\\ PG 1248$+$401        & 1.032        &   12& 50& 48.3 &   $+$39& 51& 40&   16.3 &   $-$27.6 &  N&  QSO     \\\\ CSO 174              & 0.653        &   12& 51& 00.3 &   $+$30& 25& 42&   17.0 &   $-$26.1 &  N&  QSO     \\\\ CSO 179              & 0.782        &   12& 53& 17.5 &   $+$31& 05& 50&   17.0 &   $-$26.4 &  N&  QSO     \\\\ Q 1252$+$0200        & 0.345        &   12& 55& 19.7 &   $+$01& 44& 12&   16.2 &   $-$25.3 &  Y&  QSO     \\\\ PG 1254$+$047        & 1.018        &   12& 56& 59.9 &   $+$04& 27& 34&   16.3 &   $-$27.6 &  N&  BAL QSO \\\\ PG 1259$+$593        & 0.472        &   13& 01& 12.9 &   $+$59& 02& 07&   15.9 &   $-$26.2 &  N&  QSO     \\\\ PG 1307$+$085        & 0.154        &   13& 09& 47.0 &   $+$08& 19& 49&   15.1 &   $-$24.6 &  N&  Sy 1    \\\\ PG 1309$+$355        & 0.183        &   13& 12& 17.7 &   $+$35& 15& 20&   15.6 &   $-$24.6 &  N&  Sy 1.2  \\\\ CSO 879              & 0.549        &   13& 21& 15.8 &   $+$28& 47& 20&   16.7 &   $-$25.8 &  N&  QSO     \\\\ PG 1338$+$416        & 1.204        &   13& 41& 00.8 &   $+$41& 23& 14&   16.8 &   $-$27.5 &  N&  QSO     \\\\ CSO 448              & 0.316        &   14& 24& 55.6 &   $+$42& 14& 05&   17.0 &   $-$24.3 &  Y&  QSO     \\\\ PG 1444$+$407        & 0.267        &   14& 46& 46.0 &   $+$40& 35& 06&   16.1 &   $-$24.7 &  N&  Sy 1    \\\\ PG 1522$+$101        & 1.328        &   15& 24& 24.5 &   $+$09& 58& 30&   16.2 &   $-$28.4 &  N&  QSO     \\\\ Q 1628.5$+$3808      & 1.461        &   16& 30& 13.6 &   $+$37& 58& 21&   17.7 &   $-$27.2 &  Y&  QSO     \\\\ PG 1630$+$377        & 1.478        &   16& 32& 01.1 &   $+$37& 37& 49&   16.3 &   $-$28.6 &  N&  QSO     \\\\ \\hline \\end{tabular} \\end{minipage} \\end{table*} ",
        "conclusions": "Several recent papers have used the very large numbers of quasars discovered by modern surveys to estimate their BH masses using virial approaches (e.g., Shen et al.\\ 2008; Fine et al.\\ 2008; Vestergaard \\& Osmer 2009).  The large numbers of quasars whose spectra are fit in these particular papers (between 1100 and almost 57,700) have allowed for very important new conclusions to be drawn about quasar demography.  Shen et al.\\ (2008) have found that the line widths of \\hbeta and \\mgii follow lognormal distributions with very weak dependencies on redshift and luminosity.  Their comparison of BH masses of radio-loud quasars (RLQSOs) and BALs with those of ``ordinary\" quasars (RQQSOs) shows that the mean of virial masses of RLQSOs is 0.12 dex larger than that of ordinary quasars, while that of BALs is indistinguishable from that of ordinary quasars. Fine et al.\\ (2008) used \\mgii lines to estimate BH masses, and found that the scatter in measured BH masses is luminosity dependent, showing less scatter for more luminous objects. Vestergaard \\& Osmer (2009) have studied the mass functions of BH masses at different redshifts, and found evidence for cosmic downsizing in their cosmic space density distributions.  The sample we have considered is much smaller, but each member has been carefully selected to be among the special group of RQQSOs and Seyfert galaxies already examined for optical microvariability (e.g., Carini et al.\\ 2007).  This criterion demands that modest aperture (usually 1--2 m) telescopes can make precise photometric measurements in just a few minutes, and so limits the members to the rare QSOs with bright apparent magnitudes (usually $m_V < 17.5$).  In addition, special care was taken in the fitting of the line profiles as discussed in Sections 3.1 and 3.2.  As the precision of CCD based differential photometry has typically improved to better than 0.01 mag for these  measurements made in a few minutes, the question as to the presence of microvariability in RQQSOs has been clearly settled in the affirmative (Gopal-Krishna et al.\\ 2000, 2003; Stalin et al.\\ 2004a,b; Gupta \\& Joshi 2005; Carini et al.\\ 2007; Gupta \\& Yuan 2009; Ram{\\'i}rez et al.\\ 2009).  However these papers find that the duty cycle (DC), or fraction of nights such behavior is detected, is low, roughly 10--20\\%, for RQQSOs observed for reasonably long periods (4 hours or more) during a night.  On the other hand, blazars show much more rapid variability, with DCs for such monitoring periods around 50--60\\% and when observed for $>6$ hours the blazar DCs rise to 60--85\\% (e.g., Carini 1990; Sagar et al.\\ 2004; Gupta \\& Joshi 2005).  The difference in microvariabilty amplitudes and DCs between RQQSOs and blazars can be easily understood if all the fluctuations arise from a jet very close to the central engine, but that jet only escapes that region and emits significant radio power for RL objects (e.g., Gopal-Krishna et al.\\ 2003; Gopal-Krishna, Mangalam \\& Wiita 2008). And most RLQSOs that are not blazars would only have modest (if any) Doppler boosting, so that they also have DCs comparable to those of RQQSOs, as found in the samples of Stalin et al.\\ (2005) and Ram{\\'i}rez et al.\\ (2009), or at a  level between those of RQQSOs and blazars (35--40\\%, in the sample of Gupta \\& Joshi 2005), is not surprising.  Only recently has a decent sized sample of radio intermediate quasars (RIQs), defined as those relatively rare objects with radio to optical flux ratios in the regime between those that are truly radio quiet and those normally defined as RLQSOs, been targeted for microvariability monitoring (Goyal et al.\\ 2009); their key result is that the RIQs also probably have a DC $\\leq$ 20\\%.  Therefore they are not likely to be beamed versions of RQQSOs as has been frequently suggested (Goyal et al.\\ 2009).  We conclude that there {\\bf may be a} weak negative correlation between \\hbeta EW and the Eddington ratio, $\\ell$, but there is a {\\bf significant} one between the \\mgii EW and $\\ell$ {\\bf (Fig.~\\ref{lab:lratio_ew_fwhm}; Section~5)}. This latter point has been made independently and more firmly by Dong et al.\\ (2009b) using a larger sample drawn from SDSS. We can also see from Fig.\\ 7 that there is a decline in FWHM with $\\ell$; this is unsurprising since the BH masses are proportional to $L_{edd}$.  We see from Fig.\\ 8 that there is a tendency for sources with detectable optical microvariability to have somewhat lower luminosities than those with no such detections. This is interesting, and should be confirmed by examining larger samples.  But it is not surprising, as regardless of whether the microvariability arises in accretion discs or jets, more massive BHs have naturally longer characteristic timescales that will presumably make any variations harder to detect over the course of a single night.  We also find that the BH masses estimated from the FWHMs of both the \\hbeta and \\mgii lines are reasonable, in that growth to their estimated masses from even small seed BHs is easily possible within the age of the Universe at their observed redshift if the mean $\\ell$ values are close to unity (Tables 4 and 5).  This remains true for the great majority of RQQSOs if the value of $\\ell$ we compute from the current continuum flux was constant until the time we observe them; however, this assumption does not work for 4 out of the 19 QSOs with \\hbeta lines.  It is also problematical for 1 out of the 26 QSOs with \\mgii lines (i.e., CSO 174, which is also among the difficult set of 4 using \\hbeta profiles).  We find no difference between the EWs (of both the \\hbeta or \\mgii lines) and the presence or absence of radio emission in those QSOs (Figs.\\ 4 and 7).  As discussed in the introduction, if much of the optical emission in RLQSOs comes from a jet, then we would expect the EWs of the RLQSO sources to be significantly lower than those of the RQQSOs, and if anything, they are slightly higher in our sample.  This result does not support the hypothesis (e.g., Czerny et al.\\ 2008) that RQQSOs possess jets that are producing rapid variations.  Instead it may indicate that variations involving the accretion disc (e.g., Wiita 2006) play an important role here.  Improvements to our results could be obtained through extensive searches for INOV in a larger sample of RQQSOs.  It should be useful to divide the sample of sources  based on their EWs as available from SDSS, or otherwise uniformly obtained, spectra. Such samples could be divided into three classes based on EW of \\hbeta, e.g.,  EW$<$40\\AA, 40\\AA $\\le$ EW $<$ 80\\AA\\, and EW $\\ge$ 80\\AA. Such samples should be made as homogeneous as possible on the basis of apparent magnitudes, $z$ and M$_{V}$. With such larger and homogeneous samples, the absence of a correlation between EW and DC of INOV, as found here in our modest sample, could be confirmed or shown to be unlikely."
    },
    "0910/0910.0965_arXiv.txt": {
        "abstract": "We present the results of X-ray spectral analysis of 22 active galactic nuclei (AGNs) with a small scattering fraction selected from the Second {\\it XMM-Newton} Serendipitous Source Catalogue using hardness ratios. They are candidates of buried AGNs, since a scattering fraction, which is a fraction of scattered emission by the circumnuclear photoionized gas with respect to direct emission, can be used to estimate the size of the opening part of an obscuring torus. Their X-ray spectra are modeled by a combination of a power law with a photon index of 1.5$-$2 absorbed by a column density of $\\sim$ 10$^{23-24}$ cm$^{-2}$, an unabsorbed power law, narrow Gaussian lines, and some additional soft components. We find that scattering fractions of 20 among 22 objects are less than a typical value ($\\sim$ 3\\%) for Seyfert2s observed so far. In particular, those of eight objects are smaller than 0.5\\%, which are in the range for buried AGNs found in recent hard X-ray surveys. Moreover, [\\ion{O}{3}] $\\lambda$5007 luminosities at given X-ray luminosities for some objects are smaller than those for Seyfert2s previously known. This fact could be interpreted as a smaller size of optical narrow emission line regions produced in the opening direction of the obscuring torus. These results indicate that they are strong candidates for the AGN buried in a very geometrically thick torus. ",
        "introduction": "It is widely accepted that the cosmic X-ray background (CXB) is produced by the integrated emission of faint extragalactic pointlike sources (Brandt \\& Hasinger 2005). {\\it XMM-Newton} and {\\it Chandra} resolved 80\\%$-$100\\% of the CXB at $<$ 2 keV, while the resolved fraction of the CXB at hard X-rays (8$-$12 keV) decreased to only $\\approx$ 50\\% (Worsley et al. 2005). Various observations indicate that a large fraction of active galactic nuclei (AGNs) is hidden by a large amount of cold material (e.g., Awaki et al. 1991; Comastri 2004). According to population synthesis models of the CXB (Comastri et al. 1995; Ueda et al. 2003; Gilli et al. 2007), the peak intensity of the CXB spectrum at 30 keV can be explained by considering contribution of hidden AGNs. Such a population is yet to be understood, because the direct emission from the nucleus is absorbed by surrounding cold gas and is hard to be fully explored with X-ray observations below 10 keV. Hard X-ray surveys performed with {\\it Swift}/BAT (15$-$200 keV; Markwardt et al. 2005; Tueller et al. 2008) and {\\it INTEGRAL} (10$-$100 keV; Bassani et al. 2006; Beckmann et al. 2006; Sazonov et al. 2007) are suitable for unveiling such a type of AGNs with much less selection biases than surveys at lower energies. In fact, AGNs buried in a large amount of matter have been discovered by {\\it Suzaku} follow up observations of a sample selected by the {\\it Swift/}BAT survey (Ueda et al. 2007).  In a unified model of an AGN (e.g., Antonucci 1993), torus-like gas is surrounding a supermassive black hole (SMBH), and photoionized gas is created in the opening part of the torus by radiation from the nucleus. If an AGN is observed from the torus side, absorbed direct emission and scattered light by the photoionized gas will be observed in an X-ray spectrum. The fraction of scattered light to direct emission (scattering fraction) can be used to estimate the opening angle of the torus. The fractions for AGNs found by Ueda et al (2007) are extremely small ($<$ 0.5\\%), whereas a typical value is $\\sim$ 3\\% (Turner et al. 1997; Bianchi and Guainazzi 2007). Furthermore, Winter et al (2008) found a similar type of AGNs by {\\it XMM-Newton} observations of {\\it Swift/}BAT selected AGNs. They would be buried in a geometrically thick torus with a very small opening angle assuming that the scattering fraction reflects the solid angle of the opening part of the torus. In an early stage of the evolution of galaxies and their central black holes, a large amount of gas responsible for active star formation may be closely related to obscuration of the nucleus. Therefore, AGNs almost fully covered by an absorber are an important class of objects in studying evolution of AGNs and their hosts.  Moreover, they might be significant contributors to the CXB at hard X-rays. Testing a selection technique to find such AGNs and understanding the properties of the population are of significant interest for exploring these issues.   We search for buried AGNs with a scattering fraction of 0.5\\% or less using the Second {\\it XMM-Newton} Serendipitous Source Catalogue (2XMM) and the archival data of {\\it XMM-Newton}. We selected candidate sources from the catalogue using hardness ratios (HRs) and scattering fractions calculated by analyzing spectra obtained with {\\it XMM-Newton}. The selection method of candidate sources is described in Section 2. Our results of spectral analysis are presented in Section 3 and their characteristics are discussed in Section 4. Section 5 summarizes our conclusions. We adopt ({\\it H}$_{\\rm 0}$, $\\Omega_{\\rm m}$, $\\Omega_{\\rm \\lambda}$) $=$ (70 km s$^{-1}$Mpc$^{-1}$, 0.3, 0.7) throughout this paper. ",
        "conclusions": "We searched for AGNs, whose scattered emission is very weak, from the 2XMM Catalogue. In our selection procedure, we calculated HRs expected for an object with a small scattering fraction using a model consisting of absorbed and unabsorbed power laws and 22 sources were selected as candidates from the 2XMM Catalogue. Spectral analysis was conducted using the data observed with {\\it XMM-Newton} for these 22 sources. Their X-ray spectra are represented by a combination of an absorbed power law with a column density of $\\sim$ 10$^{23-24}$ cm$^{-2}$, an unabsorbed power law, a narrow Gaussian line for the Fe K emission, and some additional components. The photon indices of the power-law components for 14 objects are in a typical range of Seyfert 2s ($\\sim$ 1.5$-$2).  The photon indices for the others, which have large uncertainties in $\\Gamma$, were fixed at 1.9 in our analysis. We found that the scattering fractions for our sample (except for NGC 4507 and IC 4995) were small compared to a typical value ($\\sim$ 3\\%) of Seyfert 2s observed in the past. In addition, those of eight sources are less than 0.5\\%. If an opening angle of a torus is responsible for the scattering fraction, objects in our sample would be hidden by a geometrically thick torus with a small opening angle.  The ratios of {\\it L}$_{\\rm 0.5-2}^{\\rm int}$/{\\it L}$_{\\rm FIR}$ for about a half of our sample are in the range observed for starburst galaxies. This result indicates that thermal emission from a collisionally ionized plasma produced by starburst may contribute to the soft X-ray emission, and true values of scattering fraction for our sample would be smaller than the value calculated in this paper.  The distribution of {\\it L}$_{\\rm 2-10}^{\\rm int}$/{\\it L}$_{\\rm [O\\ III]}^{\\rm int}$ for our sample is shifted to higher values than that for other Seyfert 2s studied so far. {\\it L}$_{\\rm [O\\ III]}^{\\rm int}$ depends on the size of the NLR that is considered to be existed in the opening part of the torus. Thus, this result also indicates that the opening angle of the obscuring matter (or scattering optical depth) for our sample is smaller than those for other Seyfert 2, and there is a bias against such a type of AGN buried in a very geometrically thick torus in surveys that rely on optical emission."
    },
    "0910/0910.5547_arXiv.txt": {
        "abstract": "We describe an implementation of compressible inviscid fluid solvers with block-structured adaptive mesh refinement on Graphics Processing Units using NVIDIA's CUDA. We show that a class of high resolution shock capturing schemes can be mapped naturally on this architecture. Using the method of lines approach with the second order total variation diminishing Runge-Kutta time integration scheme, piecewise linear reconstruction, and a Harten-Lax-van Leer Riemann solver, we achieve an overall speedup of approximately 10 times faster execution on one graphics card as compared to a single core on the host computer.  We attain this speedup in uniform grid runs as well as in problems with deep AMR hierarchies. Our framework can readily be applied to more general systems of conservation laws and extended to higher order shock capturing schemes. This is shown directly by an implementation of a magneto-hydrodynamic solver and comparing its performance to the pure hydrodynamic case. Finally, we also combined our CUDA parallel scheme with MPI to make the code run on GPU clusters. Close to ideal speedup is observed on up to four GPUs. ",
        "introduction": "\\label{sec:introduction}  Graphics Processing Units (GPUs) are specialized for math-intensive highly parallel computation, thus more transistors are devoted to data processing rather than data caching and flow control like in CPU. So the potential tremendous performance of general non-graphics computations on GPUs has recently motivated a lot of research activities on general-purpose GPU (GPGPU) computing (see. e.g. Owens et al. 2007 for a review).  NVIDIA introduced the Tesla unified graphics and computing architecture in November 2006. The Tesla architecture is built around a scalable array of multithreaded streaming multiprocessors (SM). A SM consists of eight streaming processor (SP) cores. The Tesla SM uses a new processor architecture called single-instruction, multiple-thread (SIMT). The SIMT unit creates, manages and executes up to $768$ concurrent threads in hardware with zero scheduling overhead. The SM also implements barrier synchronization intrinsic with a single instruction. The fast barrier synchronization, together with lightweight thread creation and zero-overhead thread scheduling support very fine-grained parallelism allowing thousands and even millions of threads to be invoked in kernel calls to achieve highly scalable parallel programming.  High performance supercomputing has been important in modern astrophysical research since it became available. Simulations allow astronomers to perform ``experiments\" on astronomical objects, collide stars, galaxies, or model the entire visible Universe; all situations are clearly impossible to recreate in a terrestrial laboratory. Studying the formation of stars, black holes and galaxies in the Universe is particularly challenging computationally. Their formation involves the nonlinear interplay of a range of physical processes including gravity, turbulence, magnetic field, shocks, radiation, chemistry, etc. Those questions motivated the astrophysical community to develop robust and efficient fluid codes with all the relevant physics.  Studies involving astrophysical fluid dynamics in general are benefitting tremendously from using spatial and temporal adaptive mesh refinement (AMR). This is especially so in the studies of structure formation. For example, the radius of a star is 8 orders of magnitude smaller than the size of a molecular cloud. A uniform grid code is hopeless. On the other hand, the AMR technique has been demonstrated to work well in resolving the large dynamical range involved in those problems (e.g. Abel et al. 2002; Wang \\& Abel 2009).  The mapping of computational fluid algorithms to GPU however is still at an early stage of development. Harris et al. (2003) performed cloud simulations using Stam's method (Stam 1999). This method is also used by Liu et al. (2004) for 3D flow calculations. Using finite difference methods, Brandvik \\& Pulla (2008) solved uniform grid 3D Euler equations, Elsen, LeGresley \\& Darve (2008) solved 3D Euler equations on a multi-block meshes and Zink (2008) solved Einstein's equation with uniform grid. As far as we are aware of, this work is the first on mapping an adaptive mesh finite volume solver to GPU. ",
        "conclusions": "In this work we described how to map HRSC schemes for hyperbolic conservation laws to GPU using NVIDIA's CUDA. We demonstrated that our framework as applied to the equations of inviscid compressible hydrodynamics and MHD can lead to a significant speedup. Specifically, on a Quadro FX 5600 card, saw approximately a factor of ten speedup compared to a single 3 GHz CPU core.  An important purpose of our GPU parallel scheme design is achieving good scalability. In typical fluid simulations with as many as millions of cells per grid, our parallelization scheme will launch millions of threads at the same time on GPU to perform the computation. This exceeds the Quadro FX 5600's ability of running $12288$ threads concurrently by more than two orders of magnitude. Thus we expect that the speedup factor of our parallelization scheme will increase linearly with the number of SMs on the GPU. This makes our scheme highly scalable to future generation of graphics hardwares.  One important topic we would like to concentrate on in the near future is sparse matrix solvers for Poisson equation. If one can also speed up Poisson solvers by a similar factor on GPU, then coupled with the fluid solvers we implemented in this work, a whole range of astrophysical simulations will be open to processing on the GPU as we can then model astrophysical fluids with self-gravity, viscosity and other non-ideal effects. An implementation of Poisson solver will also make it possible to use projection method for divergence cleaning and implicit fluid solvers."
    },
    "0910/0910.0224_arXiv.txt": {
        "abstract": "During the last three decades, many papers have reported the existence of a luminosity--metallicity or mass--metallicity ($M$--$Z$) relation for all kinds of galaxies: The more massive galaxies are also the ones with more metal-rich interstellar medium.  We have obtained the mass-metallicity relation at different lookback times for the same set of galaxies from the Sloan Digital Sky Survey (SDSS), using the stellar metallicities estimated with our spectral synthesis code {\\sc starlight}. Using stellar metallicities has several advantages: We are free of the biases that affect the calibration of nebular metallicities; we can include in our study objects for which the nebular metallicity cannot be measured, such as AGN hosts and passive galaxies; we can probe metallicities at different epochs of a galaxy evolution.  We have found that the $M$--$Z$ relation steepens and spans a wider range in both mass and metallicity at higher redshifts for SDSS galaxies. We also have modeled the time evolution of stellar metallicity with a closed-box chemical evolution model, for galaxies of different types and masses. Our results suggest that the $M$--$Z$ relation for galaxies with present-day stellar masses down to $10^{10} M_\\odot$ is mainly driven by the star formation history and not by inflows or outflows. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.1222_arXiv.txt": {
        "abstract": "Although primarily aimed at the galactic archeology and evolution, automated all-sky spectroscopic surveys (RAVE, SDSS) are also a valuable source for the binary star research community. Identification of double-lined spectra is easy and it is not limited by the rare occurrences of eclipses. When the spectrum is properly classified, its atmospheric parameters can be calculated by comparing the spectrum with the best fit atmosphere model. We present the analysis of the binary stars from the sample of roughly 250.000 RAVE survey spectra. The classification and binary discovery method is based on the correlation function analysis. The comparison of these spectra with the model shows that it is possible to estimate the essential atmospheric parameters relatively well. Large number of such estimates and the fact that RAVE consists of a magnitude selected sample without any color cuts makes it suitable for a binary star population study. ",
        "introduction": "The standard way of solving binary star systems is through the use of sophisticated models, the Wilson-Devinney code for example. This approach is very appealing because it is possible to precisely recover a number of system's parameters. The disadvantage is in the fact that for a reliable solution one needs a large set of photometric as well as spectroscopic observations. The latter only serve as a mean for measuring radial velocities, usually discarding the abundance of other information carried by spectra. It has been proven many times on spectra of single stars \\citep[e.g.~][to name just a few]{bailer-jones1997,katz1998,decin2004,ocvirk2006,koleva2009} that it is possible to estimate some of the astrophysical parameters (e.g.~temperature, gravity, and metallicity) by fitting the observed spectrum with the stellar atmosphere model. The same can be done in case of binary stars. This approach might be particularly useful in the pipelines of spectroscopic surveys, where parameters of large amount of observed objects have to be extracted automatically.  In the following section we will briefly review the RAVE survey itself and present the results of estimation of the most significant parameters from the spectra of binary stars in the RAVE survey. Before that we will also deal with another problem that has to be addressed when working with large datasets of unknown objects, the discovery of binary stars. Because it is not known in advance if some spectrum belongs to a binary star or some other type of peculiar object, an automatic method has to be used to properly classify all spectra. ",
        "conclusions": "The described classification and parameter estimation methods gave satisfactory results to some extent, but there is still room for improvement. Other methods like the mentioned nested sampling seem very promising for such tasks. Also, the planned use of stellar evolution models will additionally constrain the errors of calculated parameters, yielding better results.  The forthcoming missions like Gaia and the Hermes survey, will provide even more data with higher resolution. The need for truly automatic processing pipelines will be even grater, so it is important to gain as much experience as possible until then."
    },
    "0910/0910.4541_arXiv.txt": {
        "abstract": "\\noindent Using a Thomas-Fermi model, we calculate the structure of the electrosphere of the quark antimatter nuggets postulated to comprise much of the dark matter.  This provides a single self-consistent density profile from ultrarelativistic densities to the nonrelativistic Boltzmann regime that use to present microscopically justified calculations of several properties of the nuggets, including their net charge, and the ratio of MeV to 511~keV emissions from electron annihilation.  We find that the calculated parameters agree with previous phenomenological estimates based on the observational supposition that the nuggets are a source of several unexplained diffuse emissions from the Galaxy.  As no phenomenological parameters are required to describe these observations, the calculation provides another nontrivial verification of the dark-matter proposal. The structure of the electrosphere is quite general and will also be valid at the surface of strange-quark stars, should they exist. ",
        "introduction": "\\label{sec:introduction}\\noindent In this paper we explore some details of a testable and well-constrained model for dark matter~\\cite{Zhitnitsky:2002qa, Oaknin:2003uv, Zhitnitsky:2006vt, Forbes:2006ba, Forbes:2008uf} in the form of quark matter as antimatter nuggets.  In particular, we focus on physics of the ``electrosphere'' surrounding these nuggets: It is from here that observable emissions emanate, allowing for the direct detection of these dark-antimatter nuggets.  We first provide a brief review of our proposal in Sec.~\\ref{sec:proposal}, then describe the structure of the electrosphere of the nuggets using a Thomas-Fermi model in Sec.~\\ref{sec:electr-struct}.  This allows us to calculate the charge of the nuggets, and to discuss how they maintain charge equilibrium with the environment.  We then apply these results to the calculation of emissions from electron annihilation in Sec.~\\ref{sec:diffuse-1-20}, computing some of the phenomenological parameters introduced in~\\cite{Lawson:2007kp} required to explain current observations.  The values computed in the present paper are consistent with these phenomenologically motivated values, providing further validation of our model for dark matter.  The present results concerning the density profile of the electrosphere may also play an important role in the study of the surface of quark stars, should they exist. ",
        "conclusions": "\\noindent Solving the relativistic Thomas-Fermi equations, we determined the charge and structure of the positron electrosphere of quark antimatter nuggets that we postulate could comprise the missing dark-matter in our Universe.  We found the structure of the electrosphere to be insensitive to the size of the nuggets, as long as they are large enough to be consistent with current terrestrial based detector limits, and hence, can make unambiguous predictions about electron annihilation processes.  To test the dark-matter postulate further, we used the structure of this electrosphere to calculate the annihilation spectrum for incident electronic matter.  The model predicts two distinct components: a 511 keV emission line from decay through a positronium intermediate and an MeV continuum emission from direct-annihilation processes deep within the electrosphere.  By fixing the general normalization to the measured 511 keV line intensity seen from the core of the Galaxy, our model makes a definite prediction about the intensity and spectrum of the MeV continuum spectrum without any additional adjustable model parameters: Our predictions are based on well-established physics.  As discussed in~\\cite{Lingenfelter:2009fk}, a difficulty with most other dark-matter explanations for the 511 keV emission is to explain the large $\\sim$100\\% observed positronium fraction -- positrons produced in hot regions of the Galaxy would produce a much smaller fraction.  Our model naturally predicts this observed ratio everywhere.  \\textit{A priori}, there is no reason to expect that the predicted MeV spectrum should correspond to observations: typically two uncorrelated emissions are separated by many orders of magnitude.  We find that the phenomenological parameter $\\chi$~(\\ref{eq:branching}) required to explain the relative normalization of MeV emissions arises naturally from our microscopic calculation.  This is highly nontrivial because it requires a delicate balance between the two annihilation processes from the semirelativistic region of densities that is sensitive to the semirelativistic self-consistent structure of the electrosphere outside the range of validity of the analytic ultrarelativistic and nonrelativistic regimes.  (See Fig.~\\ref{fig:antinugget_profiles} and~\\ref{fig:survival}).  If the predicted emission were several orders of magnitude too large, the observations would have ruled out our proposal.  If the predicted emissions were too small, the proposal would not have been ruled out, but would have been much less interesting.  Instead, we are left with the intriguing possibility that both the 511 keV spectrum and much of the MeV continuum emission arise from the annihilation of electrons on dark-antimatter nuggets.  While not a smoking gun -- at least until the density and velocity distributions of matter and dark-matter are much better understood -- this provides another highly nontrivial test of the proposal that, \\textit{a priori}, could have ruled it out.  Both the formal calculations and the resulting structure presented here -- spanning density regimes from ultrarelativistic to nonrelativistic -- are similar to those relevant to electrospheres surrounding strange-quark stars should they exist. Therefore, our results may prove useful for studying quark star physics.  In particular, problems such as bremsstrahlung emission from quark stars originally analyzed in~\\cite{Jaikumar:2004rp} (and corrected in~\\cite{Caron:2009zi}) that uses only ultrarelativistic profile functions.  The results of this work can be used to generalize the corresponding analysis for the entire range of allowed temperatures and chemical potentials. Another problem which can be analyzed using the results of the present work is the study of the emission of energetic electrons produced from the interior of quark stars. As advocated in~\\cite{Charbonneau:2009ax}, these electrons may be responsible for neutron star kicks, helical and toroidal magnetic fields, and other important properties that are observed in a number of pulsars, but are presently unexplained.  Finally, we would like to emphasize that this mechanism demonstrates that dark matter may arise from \\emph{within} the standard model at the \\QCD\\ scale,\\footnote{The axion is another dark-matter candidate arising from the \\QCD\\ scale, but with fewer observational consequences.} and that exotic new physics is not required.  Indeed, this is naturally suggested by the ``cosmic coincidence'' of almost equal amounts of dark and visible contributions to the total density $\\Omega_{\\text{tot}} = 1.011(12)$ of our Universe~\\cite{Amsler:2008zz}: \\begin{equation*} \\Omega_{\\text{dark-energy}} : \\Omega_{\\text{dark-matter}} : \\Omega_{\\text{visible}} \\approx 17:5:1. \\end{equation*} The dominant baryon contribution to the visible portion $\\Omega_{\\text{visible}} \\approx \\Omega_{B}$ has an obvious relation to \\QCD\\ through the nucleon mass $m_{\\textsc{n}} \\propto \\Lambda_{\\QCD}$ (the actual quark masses arising from the Higgs mechanism contribute only a small fraction to $m_{\\textsc{n}}$). Thus, a \\QCD\\ origin for the dark components would provide a natural solution to the extraordinary ``fine-tuning'' problem typically required by exotic high-energy physics proposals.  Our proposal here solves the matter portion of this coincidence.  For a proposal addressing the energy coincidence we refer the reader to \\cite{Urban:2009vy,*Urban:2009yg} and references therein.\\footnote{The idea concerns the anomaly that solves the famous axial $U(1)_{A}$ problem, giving rise to an $\\eta'$ mass that remains finite, even in the chiral limit. Under some plausible and testable assumptions about the topology of our Universe, the anomaly demands that the cosmological vacuum energy depend on the Hubble constant $H$ and \\QCD\\ parameters as $\\rho_{\\textsc{de}} \\sim H m_q\\langle\\bar{q}q\\rangle /m_{\\eta'} \\approx (4 \\times 10^{-3}\\text{ eV})^4$ -- tantalisingly close to the value $\\rho_{\\textsc{de}} = [1.8(1)\\times 10^{-3}\\text{ eV}]^4$ observed today~\\cite{Amsler:2008zz}.}"
    },
    "0910/0910.1478_arXiv.txt": {
        "abstract": "In this work we study the internal spatial structure of 16 open clusters in the Milky Way spanning a wide range of ages. For this, we use the minimum spanning tree method (the $Q$ parameter, which enables one to classify the star distribution as either radially or fractally clustered), King profile fitting, and the correlation dimension ($D_c$) for those clusters with fractal patterns. On average, clusters with fractal-like structure are younger than those exhibiting radial star density profiles. There is a significant correlation between $Q$ and the cluster age measured in crossing time units. For fractal clusters there is a significant correlation between the fractal dimension and age. These results support the idea that stars in new-born clusters likely follow the fractal patterns of their parent molecular clouds, and eventually evolve toward more centrally concentrated structures. However, there can exist stellar clusters as old as $\\sim 100$ Myr that have not totally destroyed their fractal structure. Finally, we have found the intriguing result that the lowest fractal dimensions obtained for the open clusters seem to be considerably smaller than the average value measured in galactic molecular cloud complexes. ",
        "introduction": "The hierarchical structure observed in some open clusters is presumably a consequence of its formation in a turbulent medium with an underlying fractal structure (\\cite[Elmegreen \\& Scalo 2004]{Elm04}). Otherwise, open clusters having central star concentrations with radial star density profiles likely reflect the dominant role of gravity, either on the primordial gas structure or as a result of a rapid evolution from a more structured state (\\cite[Lada \\& Lada 2003]{Lad03}). Therefore, the analysis of the distribution of stars may yield information on the formation process and early evolution of open clusters. It is necessary, however, that this kind of analysis is done by measuring the cluster structure in an objective, quantitative, as well as systematic way. Here we study the internal spatial structure in a sample of 16 open clusters spanning a wide range of ages. ",
        "conclusions": "Our results support the idea that stars in new-born cluster likely follow the fractal patterns of their parent molecular clouds, and that eventually evolve toward more centrally concentrated structures (see \\cite[Schemja \\& Klessen 2006]{Sch06}; \\cite[Schmeja et al. 2008]{Sch08}, \\cite[2009]{Sch09}; \\cite[S\\'anchez et al.  2007a]{San07a}, \\cite[2009]{San09}). However, this seems to be only an overall trend. The very young cluster $\\sigma$ Orionis (age $\\sim 3$ Myr) exhibits a radial density gradient with $Q \\simeq 0.88$ (\\cite[Caballero 2008]{Cab08}). On the other hand, Table~\\ref{data} shows open clusters as old as $\\sim 100$ Myr that have not totally destroyed their clumpy structure (for example, both NGC~1513 and NGC~1647 have $Q \\sim 0.7$). \\cite[Goodwin \\& Whitworth (2004)]{Goo04} simulated the dynamical evolution of young clusters and showed that the survival of the initial substructure depends strongly on the initial velocity dispersion. Fractal clusters with a low velocity dispersion tend to erase their substructure rather quickly. However, if the velocity dispersion is high, such that the cluster remains supported against its own gravity or even expands, then significant levels of substructure can survive for several crossing times. Thus, our results give some observational support to \\cite[Goodwin \\& Whitworth's (2004)]{Goo04} simulations.  From Fig.~\\ref{correlaDc}, we can see that clusters with the smallest correlation dimensions ($D_c = 1.74$) would have three-dimensional fractal dimensions around $D_f \\sim 2.0$ (estimated from previous papers, see e.g. Fig.~1 in \\cite[S\\'anchez \\& Alfaro 2008]{San08}). This is a very interesting result because this value is considerably smaller than the average value estimated for galactic molecular clouds in recent studies, which is $D_f \\simeq 2.6-2.7$ (\\cite[S\\'anchez et al. 2005]{San05}, \\cite[S\\'anchez et al. 2007b]{San07b}). Young, new-born stars probably will reflect the conditions of the interstellar medium from which they were formed. Therefore, a group of stars born from the same cloud, i.e. born at almost the same place and time, should have a fractal dimension similar to that of the parent cloud. If the fractal dimension of the interstellar medium has a nearly universal value around 2.6-2.7, then how can some clusters exhibit such small fractal dimension values? Perhaps some clusters may develop some kind of substructure starting from an initially more homogeneous state. This possibility has been confirmed in numerical simulations (\\cite[Goodwin \\& Whitworth 2004]{Goo04}), although some coherence in the initial velocity dispersion is required. Another explanation is that this difference is a consequence of a more clustered distribution of the densest gas from which stars form at the smallest spatial scales in the molecular cloud complexes, according to a multifractal scenario (\\cite[Chappell \\& Scalo 2001]{Cha01}). Maybe the star formation process itself modifies in some (unknown) way the underlying geometry generating distributions of stars that can be very different from the distribution of gas in the parental clouds. Finally, one possibility is that the fractal dimension of the interstellar medium in the Galaxy does not have a universal value and therefore some regions form stars distributed following more clustered patterns. There is no a priori reason for assuming that $D_f$ has nearly the same value everywhere in the Galaxy independently of either the dominant physical processes or environmental variables. Recent simulations of supersonic isothermal turbulence done by \\cite[Federrath et al. (2009)]{Fed09} showed that compressive forcing yields fractal dimension values for the interstellar medium significantly smaller ($D_f \\sim 2.3$) compared to solenoidal forcing ($D_f \\sim 2.6$). Thus, $D_f$ could be very different from region to region in the Galaxy depending on the main physical processes driving the turbulence. At least at galactic scales, it has been shown that there are significant differences in the fractal dimension of the distribution of star forming sites among the galaxies, contrary to the universal picture previously claimed in the literature (\\cite[see S\\'anchez \\& Alfaro 2008]{San08}). So that the possibility of a non-universal fractal dimension for the interstellar medium in the Galaxy cannot, in principle, be ruled out.  \\begin{figure}[t] \\begin{minipage}[t]{0.48\\linewidth} \\includegraphics[scale=0.52]{correlaQ.eps} \\caption{Structure parameter $Q$ as a function of the logarithm of age divided by the tidal radius, which is nearly proportional to age in crossing times units. The dashed line at $Q=0.8$ roughly separates radial from fractal clustering. The best linear fit (equation at the top) is represented by a solid line.} \\label{correlaQ} \\end{minipage} \\hfill \\begin{minipage}[t]{0.48\\linewidth} \\includegraphics[scale=0.52]{correlaDc.eps} \\caption{Calculated correlation dimension as a function of age (in crossing time units). The best linear fit (equation at the top) is represented by a solid line. As reference, horizontal dashed lines indicate the values corresponding to three-dimensional distributions with fractal dimensions of $D_f=2.0$ and $2.5$.} \\label{correlaDc} \\end{minipage} \\end{figure}"
    },
    "0910/0910.4294_arXiv.txt": {
        "abstract": "{Although the {{\\it internal linear combination}} method (ILC) is a technique widely used for the separation of the Cosmic Microwave Background signal from the Galactic foregrounds, its characteristics are not yet well defined. This can lead to misleading conclusions about the actual potentialities and limits of such approach in real applications. Here we discuss briefly some facts about ILC that to our knowledge are not fully worked out in literature and yet have deep effects in the interpretation of the results.} ",
        "introduction": "A widely used approach for the separation of the {\\it cosmic microwave background} (CMB) from the diffuse Galactic background is the {\\it internal linear combination} method (ILC). For instance, this method was adopted in the reduction of the data from the Wilkinson Microwave Anisotropy Probe (WMAP) satellite for CMB observations \\citep{ben03}. Its success is due to the fact that, among the separation techniques, ILC calls for the smallest number of {\\it a priori} assumptions. If the data are in the form of $N_o$ maps, taken at different frequencies and containing $N_p$ pixels each, the model on which ILC is based is \\begin{equation} \\label{eq:image} \\Si = \\Scmb + \\Sgal + \\Nmthi(p). \\end{equation} Here, $\\Si$ provides the value of the $p$th pixel for a map obtained at channel ``$~i~$'' \\footnote{In the present work, $p$ indexes pixels in the classic spatial domain. However, the same formalism applies if other domains are considered, for example, the Fourier one.}, whereas $\\Scmb$, $\\Sgal$ and $\\Nmthi(p)$ are the contributions due to the CMB, the diffuse Galactic foreground and the experimental noise, respectively. Although not necessary, often it is assumed that all of these contributions are representable by means of stationary random fields. Moreover, without loss of generality, for ease of notation the random fields are supposed as the realization of zero-mean spatial processes. The basic idea behind model~(\\ref{eq:image}) is that, contrary to the components that form the Galactic background, CMB is independent of the observing channel. ILC exploits this fact averaging $N_o$ images $\\{ \\Si \\}_{i=1}^{N_o}$ and giving a specific weight $w_i$ to each of them so as to minimize the impact of the foreground and noise \\citep{ben03}. This means to look for a solution of type \\begin{equation} \\label{eq:sum} \\Scmbh = \\sum_{i=1}^{N_o} w_i \\Si. \\end{equation} If the constraint $\\sum_{i=1}^{N_o} w_i = 1$ is imposed, Eq.~(\\ref{eq:sum}) becomes \\begin{equation} \\label{eq:wls} \\Scmbh = \\Scmb + \\sum_{i=1}^{N_o}  w_i [ \\Sgal + \\Nmthi(p) ]. \\end{equation} Now, from this equation it is clear that, for a given pixel ``$p$'',  the only variable terms are in the sum. Hence, under the assumption of independence of $\\Scmb$ from $\\Sgal$ and $\\Nmthi(p)$, the weights $\\{ w_i \\}$ have to minimize the variance of $\\Scmbh$, i.e. \\begin{align} & \\{ w_i \\} = \\underset{\\{ w_i \\} }{\\arg\\min} \\nonumber \\\\ & {\\rm VAR} \\left[\\Scmb \\right] + {\\rm VAR}\\left[\\sum_{i=1}^{N_o}  w_i (\\Sgal + \\Nmthi(p)) \\right], \\end{align} where ${\\rm VAR}[s(p)]$ is the {\\it expected variance} of $s(p)$. If $\\Sib$ denotes a {\\bf row vector} such as $\\Sib = [S^{(i)}(1), S^{(i)}(2), \\ldots, S^{(i)}(N_p)]$ and the $N_o \\times N_p$ matrix $\\Sb$ is defined as \\begin{equation} \\Sb = \\left( \\begin{array}{c} \\Sb^{(1)} \\\\ \\Sb^{(2)} \\\\ \\vdots \\\\ \\Sb^{(N_o)} \\end{array} \\right), \\end{equation} then Eq.~(\\ref{eq:image}) becomes \\begin{equation} \\label{eq:basicmm} \\Sb = \\Scmbb + \\Sgalb + \\Nmthb. \\end{equation} In this case, the weights are given by \\citep{eri04} \\begin{equation} \\label{eq:wr} \\wb = \\frac{\\Cb_{\\Sb}^{-1} \\oneb}{\\oneb^T \\Cb_{\\Sb}^{-1} \\oneb}, \\end{equation} where $\\Cb_{\\Sb}$ is the $N_o \\times N_o$ cross-covariance matrix of the random processes that generate $\\Sb$, i.e. \\begin{equation} \\Cb_{\\Sb} = {\\rm E}[\\Sb \\Sb^T], \\end{equation} and $\\oneb = (1, 1, \\ldots, 1)^T$ is a column vector of all ones. Here, ${\\rm E}[.]$ denotes the {\\it expectation operator}. Hence, the ILC estimator takes the form \\begin{align} \\Scmbhb & = \\wb^T \\Sb, \\label{eq:wlss} \\\\ & = \\alpha \\oneb^T \\Cb_{\\Sb}^{-1} \\Sb, \\label{eq:basic} \\end{align} with $\\oneb^T \\wb = 1$ and the scalar quantity $\\alpha$ given by \\begin{equation} \\label{eq:alpha} \\alpha = [ \\oneb^T \\Cb_{\\Sb}^{-1} \\oneb]^{-1}. \\end{equation}  In practical applications, matrix $\\Cb_{\\Sb}$ is unknown and has to be estimated from the data. Typically, this is done by means of the estimator \\begin{equation} \\label{eq:C} \\widehat \\Cb_{\\Sb} = \\frac{1}{N_p} \\Sb \\Sb^T. \\end{equation} In this case, the ILC estimator is given by Eqs.(\\ref{eq:wlss})-(\\ref{eq:alpha}) with $\\Cb_{\\Sb}$ and $\\wb$ replaced, respectively, by $\\widehat \\Cb_{\\Sb}$ and \\begin{equation} \\label{eq:w} \\widehat \\wb = \\frac{\\widehat \\Cb_{\\Sb}^{-1} \\oneb}{\\oneb^T \\widehat \\Cb_{\\Sb}^{-1} \\oneb}. \\end{equation} ",
        "conclusions": "From the analysis presented above, it is evident that the results provided by ILC have to be interpreted with extreme caution. This techniques suffers many drawbacks that, if not properly taken into account, can lead to misleading if not wrong conclusions. In particular, ILC should be used only in situations where matrix $\\widehat\\Cb_{\\Sb}$ is (close to be) singular, i.e. only when one can be certain that model~(\\ref{eq:model})-(\\ref{eq:Amatrix}) holds with $N_o > N_c+1$. In the contrary case, the separation operated through ILC has to be expected inaccurate."
    },
    "0910/0910.3361_arXiv.txt": {
        "abstract": "We studied how the intergalactic magnetic field (IGMF) affects the propagation of super-GZK protons that originate from extragalactic sources within the local GZK sphere. Toward this end, we set up hypothetical sources of ultra-high-energy cosmic-rays (UHECRs), virtual observers, and the magnetized cosmic web in a model universe constructed from cosmological structure formation simulations. We then arranged a set of reference objects mimicking active galactic nuclei (AGNs) in the local universe, with which correlations of simulated UHECR events are analyzed. With our model IGMF, the deflection angle between the arrival direction of super-GZK protons and the sky position of their actual sources is quite large with the mean value of $\\langle \\theta \\rangle \\sim 15^{\\circ}$ and the median value of $\\tilde \\theta \\sim 7 - 10^{\\circ}$. On the other hand, the separation angle between the arrival direction and the sky position of nearest reference objects is substantially smaller with $\\langle S \\rangle \\sim 3.5 - 4^{\\circ}$, which is similar to the mean angular distance in the sky to nearest neighbors among the reference objects. This is a direct consequence of our model that the sources, observers, reference objects, and the IGMF all trace the matter distribution of the universe. The result implies that extragalactic objects lying closest to the arrival direction of UHECRs are not necessary their actual sources. With our model for the distribution of reference objects, the fraction of super-GZK proton events, whose closest AGNs are true sources, is less than 1/3. We discussed implications of our findings for correlation studies of real UHECR events. ",
        "introduction": "The nature and origin of ultra-high-energy cosmic rays (UHECRs), especially above the so-called Greisen-Zatsepin-Kuz'min (GZK) energy of $E_{\\rm GZK}\\approx$ 50 EeV (1 EeV $=10^{18}$ eV), has been one of most perplexing puzzles in astrophysics over five decades and still remains to be understood \\citep[see][for a review]{nw00}. The highest energy CR detected so far is the Fly's Eye event with an estimated energy of $\\sim 300$ EeV \\citep{flyseye94}. At these high energies protons and nuclei cannot be confined and accelerated effectively within our Galaxy, so the sources of UHECRs are likely to be extragalactic.  At energies higher than $E_{\\rm GZK}$, it is expected that protons lose energy and nuclei are photo-disintegrated via the interactions with the cosmic microwave background radiation (CMB) along their trajectories in the intergalactic space \\citep{grei66,zk66,psb76}. The former is known as the GZK effect. So a significant suppression in the energy spectrum above $E_{\\rm GZK}$ could be regarded as an observational evidence for the extragalactic origin of UHECRs \\citep[see, e.g.,][]{bgg06}. However, the accurate measurement of the UHECR spectrum is very difficult, partly because of extremely low flux of UHECRs. But a more serious hurdle is the uncertainties in the energy calibration inherent in detecting and modeling extensive airshower events \\citep[e.g.,][]{nw00, watson06}. Nevertheless, both the Yakutsk Extensive Air Shower Array (Yakutsk) and the High Resolution Fly's Eyes (HiRes) reported observations of the GZK suppression \\citep{yakutsk04,hires08a}, while the Akeno Giant Air Shower Array (AGASA) claimed a conflicting finding of no suppression \\citep{agasa04}. More recent data from the Pierre Auger Observatory (Auger) support the existence of the GZK suppression \\citep{auger08b,auger09c}. Below $E_{\\rm GZK}$, however, the four experiments reported the fluxes that are different from each other by up to a factor of several, implying the possible existence of systematic errors in their energy calibrations \\citep{bere09}.  The overall sky distribution of the arrival directions of UHECRs below $E_{GZK}$ seems to support the isotropy hypothesis \\citep[see, e.g.,][]{nw00}. This is consistent with the expectation of uniform distribution of extragalactic sources; the interaction length (i.e. horizon distance) of protons below $E_{GZK}$ is a few Gpc and the universe can be considered homogeneous and isotropic on such a large scale. The horizon distance for super-GZK events, however, decreases sharply with energy and $R_{\\rm GZK} \\sim 100 $ Mpc for $E=100$EeV \\citep{bg88}. The matter distribution inside the local GZK horizon ($R_{\\rm GZK}$) is inhomogeneous. Since powerful astronomical objects are likely to form at deep gravitational potential wells, we expect the distribution of the UHECR sources would be inhomogeneous as well. Hence, if super-GZK proton events point their sources, their arrival directions should be anisotropic.  The anisotropy of super-GZK events, hence, has been regarded to provide an important clue that unveils the sources of UHECRs. So far, however, the claims derived from analyses of different experiments are often tantalizing and sometimes conflicting. For instance, with an excessive number of pairs and one triplet in the arrival direction of CRs above 40 EeV, the AGASA data support the existence of small scale clustering \\citep{agasa96,agasa99}. On the other hand, the HiRes stereo data are consistent with the hypothesis of null clustering \\citep{hires04,hires09}. The auto-correlation analysis of the Auger data reported a weak excess of pairs for $E > 57$ EeV \\citep{auger08a}. In addition, the Auger Collaboration found a correlation between highest energy events and the large scale structure (LSS) of the universe using nearby active galactic nuclei (AGNs) in the \\citet{vcv06} catalog \\citep{auger07a,auger08a,auger09b} as well as using nearby objects in different catalogs \\citep{auger09a}. A correlation between highest AGASA events with nearby galaxies from SDSS was reported \\citep{tns09}. The HiRes data, however, do not show such correlation of highest energy events with nearby AGNs \\citep{hires08b}, but instead show a correlation with distant BL Lac objects \\citep{hires06}.  The interpretation of anisotropy and correlation analyses is, however, complicated owing to the intervening galactic magnetic field (GMF) and intergalactic magnetic field (IGMF); the trajectories of UHECRs are deflected by the magnetic fields as they propagate through the space between sources and us, and hence, their arrival directions are altered. Even with considerable observational and theoretical efforts, however, the nature of the GMF and the IGMF is still poorly constrained. Yet, models for the GMF generally assume a strength of $\\sim$ a few $\\mu$G and a coherence length of $\\sim 1$ kpc for the field in the Galactic halo \\citep[see, e.g.,][]{stan97}, and predict the deflection of UHE protons due to the GMF to be $\\theta \\sim$ a few degrees \\citep[see, e.g.,][]{ts08}. The situation for the IGMF in the LSS has been confusing. Adopting a model for the IGMF with the average strength of $\\langle B \\rangle \\sim 100$ nG in filaments, \\citet{sme03} showed that the deflection of UHECRs due to the IGMF could be very large, e.g., $\\theta > 20^{\\circ}$ for protons above 100 EeV. On the other hand, \\citet{dgst05} adopted a model with $\\langle B \\rangle \\sim 0.1$ nG in filaments and showed that the deflection should be negligible, e.g., $\\theta \\ll 1^{\\circ}$ for protons with 100 EeV.  Recently, \\citet{rkcd08} proposed a physically motivated model for the IGMF, in which a part of the gravitational energy released during structure formation is transferred to the magnetic field energy as a result of the turbulent dynamo amplification of weak seed fields in the LSS of the universe. In the model, the IGMF follows largely the matter distribution in the cosmic web, and the strength and coherence length are predicted to be $\\langle B \\rangle \\sim 10$ nG and $\\sim 1$ Mpc for the field in filaments. Such field in filaments is expected to induce the Faraday rotation \\citep{cr09}, which is consistent with observation \\citep{xkhd06}. With this model IGMF, \\citet{dkrc08} (Paper I hereafter) calculated the trajectories of UHE protons ($E>10$ EeV) that were injected at extragalactic sources associated with the LSS in a simulated model universe. We then estimated that only $\\sim 35$ \\% of UHE protons above 60 EeV would arrive at us with $\\theta \\leq 5^{\\circ}$ and the average value of deflection angle would be $\\langle \\theta \\rangle \\sim 15^{\\circ}$. Note that the deflection angle of $\\langle \\theta \\rangle \\sim 15^{\\circ}$ is much larger than the angular window of $3.1^{\\circ}$ used by the Auger collaboration in the study of the correlation between highest energy UHECR events and nearby AGNs \\citep{auger07a,auger08a,auger09b}.  In this contribution, as a follow-up work of Paper I, we investigate the effects of the IGMF on the arrival direction of super-GZK protons above 60 EeV coming from sources within 75 Mpc. The limiting parameters for energy and source distance are chosen to match the recent analysis of the Auger collaboration. Without knowing the true sources of UHECRs, the statistics that can be obtained with observational data from experiments are limited; some statistics that are essential to reveal the nature of sources are difficult or even impossible to be constructed. On the other hand, with data from simulations, any statistics can be explored. In that sense, simulations complement experiments. Here, with the IGMF suggested by \\citet{rkcd08}, we argue that the large deflection angle of super-GZK protons due to the IGMF is not inconsistent with the anisotropy and correlation recently reported by the Auger collaboration. However, the large deflection angle implies that the nearest object to a UHECR event in the sky is not necessarily its actual source. In Section 2, we describe our models for the LSS of the universe, IGMF, observers, and sources of UHECRs, reference objects for correlation study, and simulations. In Section 3, we present the results, followed by a summary and discussion in Section 4. ",
        "conclusions": "In the search for the nature and origin of UHECRs, understanding the propagation of charged particles through the magnetized LSS of the universe is important. At present, the details of the IGMF are still uncertain, mainly due to limited available information from observation. Here, we adopted a realistic model universe that was described by simulations of cosmological structure formation; our simulated universe represents the LSS, which is dominated by the cosmic web of filaments interconnecting clusters and groups. The distribution of the IGMF in the LSS of the universe was obtained with a physically motivated model based on turbulence dynamo \\citep{rkcd08}.  To investigate the effects of the IGMF on the arrival direction of UHECRs, we further adopted the following models. Virtual observers of about 1400 were placed at groups of galaxies, which represent statistically the Local Group in the simulated model universe. Then, we set up a set of about 500 AGN-like ``reference objects'' within 75 Mpc from each observer, at clusters of galaxies (deep gravitational potential wells) along the LSS. They represent a class of astronomical objects with which we performed a correlation analysis for simulated UHECR events. We considered three models, in which subsets of the reference objects were selected as AGN-like sources of UHECRs (see Table 1). UHE protons of $E \\ge 60$ EeV with power-low energy spectrum were injected at those sources, and the trajectories of UHE protons in the magnetized cosmic web were followed. At observer locations, the events with $E \\ge 60$ EeV from sources within a sphere of radius 75 Mpc were recorded and analyzed.  To characterize the clustering of the reference objects, we calculated the angular distance, $Q$, from a given reference object to its nearest neighbor. The mean value for our model AGNs in the simulated universe is $\\langle Q_{\\rm AGN} \\rangle = 3.68^{\\circ} \\pm 1.66^{\\circ}$, while that for 442 AGNs from the VCV catalog is $\\langle Q_{\\rm VCV} \\rangle = 3.55^{\\circ}$. This demonstrates that the two samples have a similar degree of clustering and are highly structured (e.g. $\\langle Q_{\\rm iso}\\rangle \\approx 11^{\\circ}$ for the isotropic distribution).  With our model IGMF, the deflection angle, $\\theta$, between the arrival direction of UHE protons and the sky position of their actual sources, is quite large with the mean $\\langle \\theta \\rangle \\sim 14 - 17.5^{\\circ}$ and the median $\\tilde \\theta \\sim 7 - 10^{\\circ}$, depending on models with different numbers of sources (see Table 1). On the other hand, the separation angle between the arrival direction and the sky position of nearest reference objects is substantially smaller with the mean $\\langle S_{\\rm sim} \\rangle \\sim 3.5 - 4^{\\circ}$ and the median $\\tilde S_{\\rm sim} = 2.8 - 3.5^{\\circ}$. That is, we found that while $\\langle \\theta \\rangle \\sim 4 \\langle S_{\\rm sim} \\rangle$, $\\langle S_{\\rm sim} \\rangle$ is similar to  $\\langle Q_{\\rm AGN} \\rangle$. For the Auger events of highest energies in \\citet{auger08a}, with 442 nearby AGNs from the VCV catalog as the reference objects, the mean separation angle is $\\langle S_{\\rm Auger} \\rangle = 3.23^{\\circ}$ for the 26 events,  excluding one event with large $S$ ($ \\approx 27^{\\circ}$), while $\\langle S_{\\rm Auger} \\rangle =4.13^{\\circ}$ for all the 27 events. Hence, $\\langle S_{\\rm Auger} \\rangle \\sim \\langle Q_{\\rm VCV} \\rangle$ $\\sim \\langle S_{\\rm sim} \\rangle \\sim \\langle Q_{\\rm AGN} \\rangle $. This implies that the separation angle from the Auger data would be determined primarily by the distribution of reference objects (AGNs), and may not represent the true deflection angle.  We further tested whether the distributions of separation angle, $S$, for our simulated events and for the Auger events are statistically comparable to each other. According the Kolmogorov-Smirnov test for the cumulative fraction of events, $F(\\le \\log S)$, versus $\\log S$, the significance level of the null hypothesis that the two distributions are drawn from the identical population is as large as $P\\sim 0.37$ for Model A (see Table 1). Thus, we argued that our simulation data, especially in Model A, are in a fair agreement with the Auger data. This test also showed that the model with more sources (Model A) is preferred over the models with fewer sources (Models B and C).  The fact that $\\langle \\theta \\rangle \\gg \\langle S_{\\rm sim} \\rangle$ implies that the AGNs found closest to the direction of UHECRs may not be the true sources of UHECRs. We estimated the probability of finding the true sources of UHECRs, when nearest reference objects are blindly chosen: $f(S)$ is the ratio of the number of true source identifications to the total number of simulated events. This probability is $\\sim 50 - 30$ \\% at $S \\sim 2^{\\circ}$, but decreases to $\\sim 10$ \\% at larger separation angle. On average, in less than 1 out of 3 events, the true sources of UHECRs can be identified in our simulations, when nearest reference objects are chosen.  The distribution of $\\theta$ versus $D_{\\theta}$ shows a bimodal pattern in which  $\\theta$ is on average larger either for nearby sources (for $D_{\\theta} \\la 15$ Mpc) or for distant sources (for $D_{\\theta} \\ga 30$ Mpc) with the minimum at the intermediate distance of $D_{\\theta,{\\rm min}} \\sim 20 - 30$ Mpc. The distribution of $S$ versus $D_s$ shows a similar, but weaker sign of the bimodal pattern. This behavior is a characteristic signature of the magnetized cosmic web of the universe, where filaments are the most dominant structure. When a large number of super-GZK events are accumulated, we may find the signature of the cosmic web of filaments in the $S$ versus $D_s$ distribution.  Finally, we address the limitations of our work. (1) We worked in a simulated universe with specific models for the elements such as the IGMF, observers, sources, and reference objects, but not in the real universe. So we could make only statistical statements. (2) It has been shown previously that adopting different models for the IGMF, very different deflection angles are obtained \\citep[see][]{sme03,dgst05,dkrc08}. We argue that our model for the IGMF is most plausible, since it is a physically motivated model based on turbulence dynamo without involving an arbitrary normalization \\citep{rkcd08}. Nevertheless, our IGMF model should be confirmed further by observation. (3) The sources of UHECRs may not be objects like AGNs, but could be objects extinguished a while ago, such as gamma-ray bursts \\citep[see, e.g.,][]{viet95,waxm95}, or sources spread over space like cosmological shocks \\citep[see, e.g.,][]{krj96,krb97}. The injection energy spectrum of power-law with cut-off at an arbitrary maximum energy (see Section 2.6) would be unrealistic. The IGMF in the Local Group (see Paper I), although currently little is known,  might be strong enough to substantially deflect the trajectories of UHECRs. All of those will have effects on the quantitative results, which should be investigated further. (4) Recently, the Auger collaboration disclosed the analysis, which suggests a substantial fraction of highest energy UHECRs might be iron nuclei \\citep{auger07b,auger09d}. This is in contradiction with the analysis of the HiRes data, which indicates highest energy UHECRs would be mostly protons \\citep{st07}. The issue of composition still needs to be settled down among experiments. Iron nuclei, on the way from sources to us, suffer much larger deflection than protons. Hence, if a substantial fraction of UHECRs is iron, some of our findings will change, a question which should be investigated in the future."
    },
    "0910/0910.3682_arXiv.txt": {
        "abstract": "Text{The giant planet atmospheres exhibit alternating prograde (eastward) and retrograde (westward) jets of different speeds and widths, with an equatorial jet that is prograde on Jupiter and Saturn and retrograde on Uranus and Neptune. The jets are variously thought to be driven by differential radiative heating of the upper atmosphere or by intrinsic heat fluxes emanating from the deep interior.  But existing models cannot account for the different flow configurations on the giant planets in an energetically consistent manner. Here a three-dimensional general circulation model is used to show that the different flow configurations can be reproduced by mechanisms universal across the giant planets if differences in their radiative heating and intrinsic heat fluxes are taken into account. Whether the equatorial jet is prograde or retrograde depends on whether the deep intrinsic heat fluxes are strong enough that convection penetrates into the upper troposphere and generates strong equatorial Rossby waves there. Prograde equatorial jets result if convective Rossby wave generation is strong and low-latitude angular momentum flux divergence owing to baroclinic eddies generated off the equator is sufficiently weak (Jupiter and Saturn). Retrograde equatorial jets result if either convective Rossby wave generation is weak or absent (Uranus) or low-latitude angular momentum flux divergence owing to baroclinic eddies is sufficiently strong (Neptune). The different speeds and widths of the off-equatorial jets depend, among other factors, on the differential radiative heating of the atmosphere and the altitude of the jets, which are vertically sheared. The simulations have closed energy and angular momentum balances that are consistent with observations of the giant planets. They exhibit temperature structures closely resembling those observed, and make predictions about as-yet unobserved aspects of flow and temperature structures.} ",
        "introduction": "Among the most striking features of the giant planets are the alternating zonal jets. As shown in Fig.~\\ref{fig_flow_lat}, Jupiter and Saturn have prograde equatorial jets (superrotation) that peak at $\\about 100\\,\\mathrm{m \\, s^{-1}}$ and $\\about 200$--$400\\,\\mathrm{m \\, s^{-1}}$, depending on the vertical level considered. Uranus and Neptune have retrograde equatorial jets (subrotation) that peak at $\\about 100\\,\\mathrm{m \\, s^{-1}}$ and $\\about 150$--$400\\,\\mathrm{m \\, s^{-1}}$.  Jupiter and Saturn have multiple off-equatorial jets in each hemisphere; Uranus and Neptune have only a single off-equatorial jet in each hemisphere.  Despite decades of study with a variety of flow models, it has remained obscure how these different flow configurations come about \\citep{Vasavada05}.  \\begin{figure*}[!htb] \\centering \\includegraphics{figure1} \\caption{Mean zonal velocities in the upper atmosphere of the giant planets from observations and simulations.  Jupiter: observations from the \\emph{Cassini} spacecraft \\citep{Porco03} (orange line), and in simulation at $0.75$~bar (dark blue line).  Saturn: observations from the \\emph{Voyager} spacecraft (orange line), from the \\emph{Hubble Space Telescope} (HST) \\citep{Sanchez-Lavega03} (green crosses), from the \\emph{Cassini} spacecraft at $\\about 0.06$~bar (magenta circles) and at $\\about 0.7$~bar (light blue squares) \\citep{Sanchez-Lavega07}, and in simulation at $0.1$~bar (dark blue line). Uranus: observations from the \\emph{Voyager} spacecraft (orange circles), HST (orange crosses) \\citep{Hammel01}, the Keck telescope (orange squares) \\citep{Hammel05}, and in simulation at $25.0$~mbar (dark blue line). Neptune: observations from the \\emph{Voyager} spacecraft (orange circles) and from HST (orange crosses) \\citep{Sromovsky01}, and in simulation at $25.0$~mbar (dark blue line). Differences between the statistically identical northern and southern hemispheres in the simulations are indicative of the sampling variability of the averages.} \\label{fig_flow_lat} \\end{figure*}  Existing models posit as the driver of the flow either the differential radiative heating of the upper atmosphere \\cite[e.g.,][]{Williams79,Williams03a} or the intrinsic heat fluxes emanating from the deep interior \\citep[e.g.,][]{Busse76,Heimpel05,Aurnou07,Chan08,Kaspi09}.  However, none of these models can account for the existence of equatorial superrotation on Jupiter and Saturn and equatorial subrotation on Uranus and Neptune with radiative heating, intrinsic heat fluxes, and other physical parameters consistent with observations.  For example, deep-flow models that posit intrinsic heat fluxes as the sole driver of the flow can generate equatorial superrotation, but they use heat fluxes more than $10^6$ times larger than those observed \\citep[e.g.,][]{Heimpel07}. They generate equatorial subrotation only with intrinsic heat fluxes even stronger than those for which they generate superrotation \\citep{Aurnou07}, although the intrinsic heat fluxes on the subrotating planets (Uranus and Neptune) are weaker than those on the superrotating planets (Jupiter and Saturn).  The relevance of such deep-flow models is further called into question by the eddy angular momentum fluxes they imply. Their meridional eddy fluxes of angular momentum per unit volume (taking density variations into account) have a barotropic structure: they extend roughly along cylinders concentric with the planet's spin axis over the entire depth of the fluid, typically to pressures of order $10^6$~bar \\citep[e.g.,][their Fig.~10]{Kaspi09}. But the eddy angular momentum fluxes inferred from tracking cloud features in Jupiter's and Saturn's upper tropospheres indicate that the mean conversion rate from eddy to mean-flow kinetic energy is of order $10^{-5} \\, \\mathrm{W \\ m^{-3}}$ \\citep{Ingersoll81,Salyk06,DelGenio07}. If the observed upper-tropospheric eddy fluxes of angular momentum per unit volume extended unabatedly over a layer of 50~km thickness (e.g., from about 0.3 to 2.5~bar pressure on Jupiter, or from about 0.3 to 0.9~bar pressure on Saturn), and if vertical zonal flow variations over this layer are not dramatic, the total energy conversion rate would be about $0.5\\,\\mathrm{W \\, m^{-2}}$. This is already $\\about 4\\%$ of the total energy uptake of the atmosphere from intrinsic heat fluxes and absorption of solar radiation for Jupiter, or $\\about 11\\%$ for Saturn. But the limited thermodynamic efficiency of atmospheres allows only a fraction of the total atmospheric energy uptake to be used to generate eddy kinetic energy \\citep{Lorenz55,Peixoto92}. The observations of Jupiter and Saturn therefore imply that eddy angular momentum fluxes cannot extend unabatedly over great depths and must have a baroclinic structure. Barotropic eddy angular momentum fluxes that extend to depths of order $10^6$~bar, with upper-atmospheric fluxes of similar scale and magnitude as those observed, are only possible in deep-flow models if the driving heat fluxes are several orders of magnitude greater than observed.  Similarly, shallow-flow models that posit differential radiative heating as the driver of the flow can generate equatorial superrotation, but they require artifices such as additional equatorial heat or wave sources that have no clear physical interpretation \\citep[e.g.,][]{Williams03a, Williams03b, Yamazaki05, Lian08}. It is unclear in those and other shallow-flow models \\citep[e.g.,][]{Scott08} what physical characteristics distinguish the superrotating planets from the subrotating planets.\\footnote{\\citet{Lian10} claim that different rates of latent heat release in phase changes of water may be responsible for superrotation on Jupiter and Saturn and subrotation on Uranus and Neptune. However, they impose latent heat fluxes at the lower boundary of their model that are not consistent with the observed energetics of the planets. Similar to the simulations of \\citet{Aurnou07}, they require stronger energy (latent heat) fluxes to generate subrotation than to generate superrotation. For example, the latent heat fluxes are of order $10$--$20\\,\\mathrm{W\\,m^{-2}}$ in their Jupiter and Saturn simulations and of order $1500\\,\\mathrm{W\\,m^{-2}}$ in their Uranus/Neptune simulation (Y.~Lian, pers.\\ communication, 2010). The latter are several orders of magnitude larger than the observed intrinsic heat fluxes or absorbed radiative fluxes (Table~1), which would have to drive any latent heat fluxes (energy would be required to evaporate the condensate that falls from the upper atmosphere into deeper layers).}  In \\citet{Schneider09} (SL09 hereafter), we postulated that prograde equatorial jets on the giant planets occur when intrinsic heat fluxes are strong enough that Rossby waves generated convectively in the equatorial region transport angular momentum toward the equator. Multiple off-equatorial jets, by contrast, form as a result of baroclinic instability owing to the differential radiative heating of the upper atmosphere. We introduced a general circulation model (GCM) and demonstrated with it that the postulated mechanisms can account qualitatively for large-scale flow structures observed on Jupiter. Here we use simulations with essentially the same GCM, with closed energy and angular momentum balances that are consistent with observations, to demonstrate universal formation mechanisms of jets on all the giant planets.  We show that the different flow configurations on the giant planets can be explained through consideration of the different roles played by intrinsic heat fluxes and solar radiation in generating atmospheric waves and instabilities.  Section~\\ref{s:GCM} briefly describes the GCM. Section~\\ref{s:Simulations} shows simulation results for Jupiter, Saturn, Uranus, and Neptune. Section~\\ref{s:Mechanisms} discusses the formation mechanisms of the jets in the simulations and confirms the postulated mechanisms through control simulations. Section~\\ref{s:AM_balance} discusses what the upper-atmospheric fluid dynamics, on which we focus, imply about flows at greater depth on the giant planets. Section~\\ref{s:Conclusion} summarizes the conclusions and their relevance for available and possible future observations. ",
        "conclusions": "\\label{s:Conclusion}  We have presented the first simulations of all four giant planets with closed energy and angular momentum balances that are consistent with observations. The simulations reproduce many large-scale features of the observed flows, such as equatorial superrotation on Jupiter and Saturn and equatorial subrotation on Uranus and Neptune. They exhibit temperature structures that are broadly consistent with available observations, and they reproduce many details of the observed flows, for example, their vertical structure to the extent it is known and characteristic equatorial waves observed on Jupiter. We have demonstrated that equatorial superrotation is generated if convective Rossby wave generation is strong and low-latitude angular momentum flux divergence owing to baroclinic eddies generated off the equator is sufficiently weak (Jupiter and Saturn); equatorial subrotation results if either convective Rossby wave generation is weak or absent (Uranus) or low-latitude angular momentum flux divergence owing to baroclinic eddies is sufficiently strong (Neptune).  Current computational resources limit our ability to simulate flows at depth.  However, considerations of the angular momentum balance have shown that the zonal jets should extend---generally with shear---to the depth where drag acts on them and balances the angular momentum flux divergences and convergences in the upper troposphere.  That drag acts on the zonal flow at depth is suggested by observations of eddy angular momentum fluxes on Jupiter and Saturn, and a plausible MHD drag mechanism exists. Though quantitative aspects (e.g., jet strength and shear) may be affected by our inability to resolve the flow and drag at depth, the jet formation mechanisms we discussed are not affected by it.  We expect as-yet unobserved aspects of the flow and temperature structures to be consistent with the simulations and mechanisms we presented. For example, we predict that NASA's upcoming JUNO mission to Jupiter will find evidence of zonal flows with vertical shear similar to those in Fig.~\\ref{fig_circulation}: near the equator, a strong and deep prograde jet, and away from the equator, prograde jets that weaken and retrograde jets that weaken only slightly or strengthen with depth. The thermal and gravitational signature of such zonal flows will likely be measurable by JUNO.  \\paragraph"
    },
    "0910/0910.1364_arXiv.txt": {
        "abstract": "{} {We investigated the non-thermal hard X-ray emission in the Ophiuchus cluster of galaxies. Our aim is to characterise the physical properties of the non-thermal component and its interaction with the cosmic microwave background.} {We performed spatially resolved spectroscopy and imaging using XMM-Newton data to model the thermal emission. Combining this with INTEGRAL ISGRI data, we modelled the 0.6--140 keV band total emission in the central 7 arcmin region.} {The models that best describe both PN and ISGRI data contain a power-law component with a photon index in a range 2.2--2.5.  This component produces $\\sim$10\\% of the total flux in the 1--10 keV band. The pressure of the non-thermal electrons is $\\sim$1\\% of that of the thermal electrons. Our results support the scenario whereby a relativistic electron population, which produces the recently detected radio mini-halo in Ophiuchus, also produces the hard X-rays via inverse compton scattering of the CMB photons. The best-fit models imply a differential momentum spectrum of the relativistic electrons with a slope of 3.4--4.0 and a magnetic field strength B=0.05--0.15 $\\mu$G. The lack of evidence for a recent major merger in the Ophiuchus centre allows the possibility that the relativistic electrons are produced by turbulence or hadronic collisions.} {} ",
        "introduction": "Non-thermal hard X-ray emission has been detected in several clusters of galaxies over the past few years (see Rephaeli et al., 2008 for a recent review). Since the detections have remained at the level of a few $\\sigma$, many models still remain as valid explanations for the phenomenon. The most popular explanation is the inverse compton scattering of the cosmic microwave background photons with the relativistic electrons in the cluster (e.g. Sarazin et al., 1988). In this model, the CMB photon ends up in the hard X-ray band since its energy increases by a factor of $\\sim 10^8$ via the scattering. In primary models, the relativistic electrons are originally thermal electrons that have been accelerated by a cluster merger (e.g. Sarazin \\& Lieu, 1998) or turbulence (e.g. Brunetti et al., 2001; Brunetti et al., 2004). In secondary models, the acceleration comes from hadronic collisions (e.g. Dennison, 1980; Pfrommer \\& Ensslin, 2004).  Recently a radio mini-halo was detected in the centre of the Ophiuchus cluster with the VLA at 1.4 GHz (Govoni et al., 2009; Murgia et al., 2009). This proved the existence of a population of relativistic electrons in Ophiuchus. Ophiuchus is a hot (T$\\sim$9 keV) nearby (z=0.028) cluster located close to the Galactic centre (l $\\sim 1^{\\circ}$ , b $\\sim  9^{\\circ}$). Consistently, INTEGRAL detected excess emission over the thermal component at a 4--6$\\sigma$ level in the 20--80 keV energy band in the Ophiuchus cluster (Eckert et al., 2008). The analysis was lacking a sensitive instrument for modelling the thermal component at lower energies, but the excess was  nevertheless consistent with having a non-thermal origin.  Our aim in this work is to improve the modelling of the emission in the Ophiuchus centre utilising spatially resolved XMM-Newton spectroscopy. In this work we use the EPIC instruments PN for spectral analysis and MOS2 for imaging of the central 7 arcmin region of the Ophiuchus cluster. Because of the limitations of the spatial capability of INTEGRAL, we cannot exclude regions of complex temperature structure. Rather, we examine the spatially resolved XMM-Newton data in order to obtain an accurate model for the thermal emission to be used later with the INTEGRAL data. We also update the INTEGRAL analysis with additional data.   \\begin{figure*}[tb] \\includegraphics[width=18cm,angle=0]{Oph_lc2.ps} \\vspace{-6cm} \\caption{The histogram shows the light curve of Ophiuchus in the full FOV in 1000 ks time bins in the hard band (E $>$ 10 keV for PN, left panel and  E $>$ 9.5 keV for MOS2, right panel). The dotted horizontal lines show the adopted quiescent level. \\label{lc.fig}} \\end{figure*} ",
        "conclusions": "We analysed the central r=7 arcmin region of the Ophiuchus cluster of galaxies using data obtained with the XMM-Newton EPIC and INTEGRAL ISGRI instruments. The ISGRI data yielded a 5.7$\\sigma$ detection of excess emission in the 20--120 keV band over the thermal prediction, as modeled with PN data. Our XMM-Newton analysis confirmed the existence of a cool core in Ophiuchus. The derived very long cooling time (3 $\\times 10^{9}$ yr) as well as the lack of significant merger signatures argues against a recent major merger in the Ophiuchus centre. These features are consistent with most clusters hosting a radio mini-halo (Ferrari et al., 2008), which has also been detected in Ophiuchus (Govoni et al. 2009).  In most proposed models for the non-thermal emission in clusters the same population of relativistic electrons produces both the radio emission (via synchrotron) and hard X-ray emission (via inverse compton scattering of the cosmic microwave background photons). Our results support this scenario, since the derived photon index of 2.2--2.5 for the non-thermal component corresponds to radio index of 1.2--1.5 with is consistent with the upper limit (1.7) derived from the radio observations of Ophiuchus (P\\'eres-Torrez et al.,  2009).  The non-thermal component produces $\\sim$10\\% of the total flux in the 1--10 keV band. These models imply a differential momentum spectrum of the relativistic electrons with a slope of 3.4-4.0 and a magnetic field strength B=0.05--0.15 $\\mu$G. The pressure of the non-thermal electrons is $\\sim$ 1\\% of that of the thermal electrons, i.e. the gravitational potential of the cluster is adequate for confining such a population of non-thermal electrons."
    },
    "0910/0910.4451_arXiv.txt": {
        "abstract": "Massive neutron stars may harbor deconfined quark matter in their cores. I review some recent work on the microphysics and  the phenomenology of compact stars with cores made of quark matter. This includes the equilibrium and stability of non-rotating and rapidly rotating stars, gravitational radiation from deformations in their quark cores, neutrino radiation and dichotomy of fast and slow cooling, and pulsar radio-timing anomalies. ",
        "introduction": "\\label{sec:Intro}  Because of the quark substructure of nucleons predicted by quantum chromodynamics (QCD), nuclear matter will undergo a phase transition to quark matter if squeezed to sufficiently high densities. In the quark matter phase the ``liberated'' quarks occupy continuum states which, in the low-temperature and high-density regime, arrange themselves in a Fermi sphere. The Fermi-sphere determines the form of low-lying excitation spectrum of quark matter, which resembles that of less exotic low-temperature systems found in condensed matter (\\eg electron gas or ultracold atomic vapor) or hadronic physics (\\eg nuclear or neutron matter). In analogy to these system the attractive interaction between quarks, mediated by the gluon exchange (which is responsible for the bound state spectrum of QCD, \\eg, the nucleon and mesons)  leads to quark superconductivity and superfluidity (color superconductivity) via the Bardeen-Cooper-Schrieffer mechanism~\\cite{Bailin:1983bm}. Furthermore, under stellar conditions the pairing between the two light flavors of quarks occurs at finite isospin chemical potential, \\ie, when the Fermi surfaces of up and down quarks are shifted apart by an amount which is of the same order of magnitude as the gap in the quasiparticle spectrum. Under these conditions the pairing between the fermions persists, but the actual pairing pattern  may be significantly different from that of the BCS and is likely to involve breaking of the spatial symmetries by the condensate order parameter.  During the last decade there has been a substantial progress in understanding of pairing in two-component asymmetric superconductors. Firstly, new developments on (variations of) the Larkin-Ovchinnikov-Fulde-Ferrell (LOFF) phase revealed novel lattice structures of the order parameter~\\cite{Bowers:2002xr}. Secondly, it has been suggested that such systems may actually phase separate into normal and superconducting domains~\\cite{Caldas} or the pairing may require changes in the shapes of the Fermi surfaces~\\cite{Muther:2002mc}. Controlled experiments on two-component cold atomic vapors with mismatched Fermi surfaces show that the phase separation scenario is at work. Whether this is the case in the related systems such as the isospin asymmetric nuclear matter or deconfined quark matter is still unclear. Therefore, a major goal of astrophysics of dense matter is to find and to quantify the manifestations of dense phases in the observable properties of compact stars. Here we review some recent progress towards this goal. ",
        "conclusions": ""
    },
    "0910/0910.5727_arXiv.txt": {
        "abstract": "Late X-ray flares observed in X-ray afterglows of gamma-ray bursts (GRBs) suggest late central engine activities at a few minuets to hours after the burst. A few unambiguously confirmed cases of supernova associations with nearby long GRBs imply that an accompanying supernova-like component might be a common feature in all long GRB events. These motivate us to study the interactions of a late jet, responsible for a x-ray flare, with various components in a stellar explosion, responsible for a GRB. These components include a supernova shell-like ejecta, and a cocoon that was produced when the main jet producing the GRB itself was propagating through the progenitor star. We find that the interaction between the late jet and the supernova ejecta may produce a luminous (up to $10^{49}$ erg s$^{-1}$) thermal X-ray transient lasting for $\\sim 10$ s The interaction between the late jet and the cocoon produces synchrotron-self absorbed non-thermal emission, with the observed peak X-ray flux density from 0.001 $\\mu$Jy to 1 mJy at 1 keV and a peak optical flux density from 0.01 $\\mu$Jy to 0.1 Jy (for a redshift $z= 2$). The light curve due to the late jet - cocoon interaction has very small pulse-width-to-time ratio, $\\D t/t \\approx 0.01 - 0.5$, where $t$ is the pulse peak time since the burst trigger. Identifying these features in current and future observations would open a new frontier in the study of GRB progenitor stars. ",
        "introduction": "Long duration gamma-ray bursts (GRB) -- those lasting more than 2 s -- are thought to be produced by a relativistic outflow (or jet) with a kinetic energy $\\sim 10^{51} - 10^{52}$ erg (beaming effect corrected) when a massive star collapses at the end of its nuclear burning life cycle  (see Piran 1999, 2005; M\\'{e}sz\\'{a}ros 2002 for reviews). The massive star origin of GRBs are supported by two different observations: (i) GRBs are found to be in actively star forming galaxies (e.g., Christensen et al. 2004, Castro Cer\\'on et al. 2006) or in star (especially massive ones) forming regions of the host galaxies (e.g., Paczy\\'{n}ski 1998; Bloom, Kulkarni \\& Djorgovski 2002; Fruchter et al. 2006); (ii) for a subset of nearly a dozen of GRBs, X-ray-rich GRBs and X-ray flashes, Supernova (SNa) features -- both temporally and spacially associated with the bursts (Woosley \\& Bloom 2006 for a review) -- were detected. For four of those: GRB 980425 (e.g., Galama et al. 1998), 030329 (e.g., Hjorth et al. 2003), 021211 (Della Valle et al. 2003) and 031203 (e.g., Malesani et al., 2004), the physically associated SNe were not only photometrically but also spectroscopically confirmed. The others of the subset show a late-time ($\\sim$ 10 days) SNa-like ``bump'' in the optical afterglows, with a simultaneous strong color evolution, e.g., in GRB 980326 (Bloom et al. 1999) and 011121 (Bloom et al. 2002), consistent with the hypothesis of an underlying SNa.  The {\\it Swift} satellite has recently unveiled a ``canonical'' behavior pattern in about two-thirds of GRBs' early X-ray afterglows: a rapid decline phase lasting for $\\sim 10^2$ s is followed by a shallow decay phase lasting $\\sim 10^3 - 10^4$ s, then by a ``normal'' power-law decay phase and finally by a possible jet break (Nousek et al. 2006; O'Brien et al. 2006). In addition, X-ray flares are found in about 50\\% of all {\\it Swift} bursts; they have been discovered in all of the above phases (Burrows et al. 2005, 2007; Chincarini et al. 2007). Even long before {\\it Swift}, a late X-ray flare was detected by BeppoSAX for GRB 970508 (Piro et al. 1998).  In this work we investigate a scenario in which a late jet -- responsible for producing the X-ray flares and possibly the shallow decay phase -- interacts with other components of a long-GRB stellar explosion. These components include a SNa ejecta, if a SNa explosion is accompanying the GRB event, and a cocoon created by the passage of the main GRB jet through the star (Ramirez-Ruiz et al. 2002, Matzner 2003, Zhang et al. 2004).  In Section \\ref{sec:mul-com} we argue for the existence of the late jet and multiple components of a GRB explosion, and provide physical motivations of this work. We investigate the interactions of the late jet with the SNa ejecta in Section \\ref{sec:jet-ejecta} and with the cocoon in Section \\ref{sec:jet-cocoon} and calculate their associated emissions. The predicted emissions and their detection prospects are confronted with current observational data in Section \\ref{sec:obs}. The summary and implications are given in Section \\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  Observations of X-ray flares in many GRB afterglows suggest the existence of a late jet from a long-lived central engine of a GRB at  $\\sim 10^2$ s but possibly  $10^4 - 10^5$ s after the main GRB event. Adopting the collapsing massive star origin for long-duration GRBs, and assuming that the supernova explosion to be at approximately the same time as the GRB, we have investigated the interactions of this late jet  with the SNa ejecta  and, with a cocoon that was left behind when the main GRB jet traversed the progenitor star.   We find that late jet - SNa ejecta interaction should produce a thermal transient, lasting about $\\sim 10$ s and with a peak photon energy at a few keV, that should precede or accompany the flare. The luminosity of this transient is proportional to the kinetic luminosity of the late jet and can be as high as a few $\\times 10^{48}$ erg s$^{-1}$. The luminosity is smaller if the polar cavity created by the main GRB jet is only partially filled. This thermal transient is similar to the one associated with the breakout of the main GRB jet. Although it has a lower luminosity, its later occurrence makes it easier to detect. The observation of this signal can provide another evidence for the massive-star origin of GRBs and new information on the GRB - SNa association.  The fact that no such thermal transient was observed so far implies that the late jet - SNa ejecta interaction is suppressed. This could happen if the polar region of the ejecta was evacuated by the main GRB jet and cavity has not sufficiently refilled itself, especially for a not-too-late jet, e.g., $t_F \\sim 10^2$ s, or the cavity could be kept open by a continuous, low-power jet. A strong thermal transient signal can also be blocked by the optically thick cocoon, which would be the case if the cocoon is slowly moving. An alternative possibility is of a failed supernova scenario in which the entire stellar envelope collapses in a free fall time of a few $\\times 10^2$ s. In this case there is no supernova associated with the GRB and the late jet, if it is late enough,  does not have to cross the stellar envelope.  The late jet interaction with the cocoon would cause a flare or rebrightening, superposed on the afterglow light curves, at both the optical and X-ray bands. This flare would have a pulse-width-to-time ratio $\\D t /t < 1$ (the expected distribution of $\\D t/t$ is similar to that for X-ray flares). Depending on model parameters, we find for a burst at a redshift $z=2$ that the peak flux density at optical $f_{\\nu_{opt}}$ ranges from 0.01 $\\mu$Jy to 0.1 Jy ($V$-band apparent magnitude 29 to 11.5) and at X-rays $f_{\\nu_X}$ ranges from 0.001 $\\mu$Jy to 1 mJy. For typical parameters $f_{\\nu_{opt}} \\sim 0.1$ mJy ($V$-band magnitude $\\sim 19$) and $f_{\\nu_X} \\sim 1$ $\\mu$Jy. Observational identification of this emission would verify the existence of the cocoon produced when the GRB jet traversed the progenitor star, thus it would be another confirmation of the collapsar model for long duration GRBs (e.g., MacFadyen \\& Woosley 1999; Ramirez-Ruiz et al. 2002; Matzner 2003; Zhang et al. 2004).  The late jet - cocoon interaction might have already been detected in four GRB afterglows in which simultaneous X-ray and optical flares with $\\D t /t \\ll 1$ were observed after the prompt emission has died off (see Fig. \\ref{fig:obs-cand}). From those candidate events, one can learn about the energetics of late jet and the cocoon by utilizing the emission calculation presented in this paper. Let us consider the flare event in GRB 050904 as an example.  We find the most probable model parameters -- for this burst at $z= 6.3$ and with $t_F= 70$ s -- that produce the observed peak $f_{\\nu_{opt}}$ and $f_{\\nu_X}$ (data from Bo\\\"{e}r et al. 2006; Cusumano et al. 2007; Gou, Fox \\& M\\'{e}sz\\'{a}ros 2007) to be $E_c \\approx 10^{51}$ erg, $\\G_c \\approx 20 - 50$, $E_j \\approx 10^{52}$ erg and $\\G_j \\approx 500$. Those high energetics seem consistent with the very luminous nature of both the burst and the flare.  There are three cases in which no optical flare was detected at the time of a strong X-ray flare, even though a number of optical telescopes have been observing these bursts at the time of the X-ray flares (Fig. \\ref{fig:obs-non-cand}). This shows that not all X-ray flares are due to the late jet - cocoon interaction. However, neither a late jet nor the cocoon can be ruled out in these cases. It is possible that a low-power jet preceding the late jet with a total energy of at least $10^{49}$ erg had kept the cavity in the cocoon open, so that the late jet - cocoon interaction was suppressed.  If correct this implies a low level continuous emission from the central engine at the level of $\\sim 10^{47}\\, (t/10 \\, {\\rm s})^{-2}$ erg s$^{-1}$ lasting for $\\sim 10^{2}$ s, assuming a radiation efficiency of $\\sim 0.1$.   \\begin{figure} \\centerline{ \\includegraphics[width=6cm,angle=0]{Molinari_et_al_060418.ps} } \\centerline{ \\includegraphics[width=6cm,angle=0]{Molinari_et_al_060607a.ps} } \\centerline{ \\includegraphics[width=7cm,angle=0]{Rykoff_et_al_060904b.ps} } \\caption{The three GRBs that show prominent late X-ray flares but without simultaneous optical flare apparent in the afterglow light curve. {\\it Top}: GRB 060418; {\\it Middle}: 060607a (both from Molinari et al. 2007). {\\it Bottom}: GRB 060904b; blue triangles are BAT data extrapolated to X-ray band, magenta squares are XRT data and red circles are optical data (from Rykoff et al. 2009).} \\label{fig:obs-non-cand} \\end{figure}"
    },
    "0910/0910.3953.txt": {
        "abstract": "The Standing Accretion Shock Instability (SASI) is commonly believed to be responsible for large amplitude dipolar oscillations of the stalled shock during core collapse, potentially leading to an asymmetric supernovae explosion. The degree of asymmetry depends on the amplitude of SASI, but the nonlinear saturation mechanism has never been elucidated. We investigate the role of parasitic instabilities as a possible cause of nonlinear SASI saturation.  As the shock oscillations create both vorticity and entropy gradients, we show that both Kelvin-Helmholtz and Rayleigh-Taylor types of instabilities are able to grow on a SASI mode if its amplitude is large enough. We obtain simple estimates of their growth rates, taking into account the effects of advection and entropy stratification. In the context of the advective-acoustic cycle, we use numerical simulations to demonstrate how the acoustic feedback can be decreased if a parasitic instability distorts the advected structure. The amplitude of the shock deformation is estimated analytically in this scenario. When applied to the set up of \\cite{fernandez09a}, this saturation mechanism is able to explain the dramatic decrease of the SASI power when both the nuclear dissociation energy and the cooling rate are varied. Our results open new perspectives for anticipating the effect, on the SASI amplitude, of the physical ingredients involved in the modeling of the collapsing star. ",
        "introduction": "Despite decades of active research \\citep{colgate66,bethe85}, the core collapse supernovae mechanism remains elusive. The failure of the most sophisticated 1D models to explode the majority of massive progenitors \\citep{liebendorfer01} suggests that multidimensional effects are essential for a successful explosion. Understanding the hydrodynamical instabilities responsible for this symmetry breaking, and more specifically their nonlinear dynamics, is therefore required to understand the explosion mechanism.  The region between the neutrinosphere and the shock deserves particular attention because two instabilities take place there: neutrino-driven convection \\citep{herant92,herant94,burrows95,janka96,foglizzo06}, and the newly discovered Standing Accretion Shock Instability (SASI) \\citep{blondin03, ohnishi06, foglizzo07, scheck08}. 2D simulations suggest that the complex fluid motions triggered by these instabilities could lead to a successful explosion either by helping the classical neutrino-driven mechanism \\citep{buras06b, marek09, murphy08b} or by a new mechanism based on the emission of acoustic waves from the proto-neutron star (\\cite{burrows06, burrows07a}, see however \\cite{weinberg08}). The large scale ($l=1-2$) induced asymmetry could also explain the high kick velocities of newly formed neutron stars \\citep{scheck04, scheck06} and may affect significantly their spin \\citep{blondin07a, yamasaki08}.  The linear phase of the two instabilities has been described in details by \\cite{foglizzo06, blondin06, foglizzo07,yamasaki07,fernandez09a}. Neutrino-driven convective modes with a large angular scale can be stabilized by a fast advection of matter through the gain region, whereas SASI is always dominated by large scale modes. This linear argument favors SASI as the cause of the prominent $ l=1-2 $ shock oscillations observed in the simulations. However, a theoretical understanding of the nonlinear development and saturation of SASI is still missing. This was highlighted by the unexpected dramatic decrease of the SASI power observed by \\cite{fernandez09a} when both the nuclear dissociation energy and the cooling rate are varied. To shed more light on this issue, this paper proposes a predictive saturation mechanism for SASI. This is a first step toward understanding how the amplitude of SASI depends on the physical ingredients of the model (nuclear dissociation, equation of state, heating rate, rotation, magnetic fields).  We propose that the saturation takes place when a parasitic instability (also called secondary instability) grows on the dominant SASI mode. These parasites feed upon its energy and destroy its coherence, leading to a turbulent flow and the saturation of the growing SASI mode. This saturation mechanism by parasitic instabilities is similar to the one proposed for the saturation of the MRI by \\cite{Goodman94, Pessah09a}.  Two types of instabilities are considered as potential parasites for the SASI mode: the Kelvin-Helmholtz instability (hereafter KHi) grows near a maximum of vorticity, and the Rayleigh-Taylor instability (hereafter RTi) grows on negative entropy gradients. In either case the cause of the instability lies within the SASI mode, and therefore the growth rate of these parasites increases with the amplitude of the shock oscillations. The parasites are thus able to affect significantly the dynamics of the flow only if the amplitude of SASI oscillations exceeds a certain threshold. The main objective of this work is to estimate this threshold.  Understanding the saturation mechanism of SASI requires in principle some understanding of the mechanism underlying its linear growth. The saturation scenario we propose is not restricted to a particular mechanism for the growth of SASI, but it can be understood more precisely in the framework of the advective-acoustic cycle \\citep{foglizzo02, blondin03, ohnishi06, foglizzo07, scheck08, fernandez09a, foglizzo09}. In this mechanism, a shock deformation creates an entropy-vorticity wave, whose downward advection generates an acoustic feedback. This acoustic wave further deforms the shock, thus closing the unstable cycle. If the coherence of the advected wave were broken by a parasitic instability before it created the acoustic feedback, this cycle would be stabilized and SASI saturated.  In Sect.~2, we explain our method and approximations. Sections 3 and 4 are devoted to the linear study of the KHi and RTi in our setup. In Sect.~5 we study how the development of the parasites breaks the coherence of a SASI mode and decreases of the acoustic feedback. In Sect.~6 we apply our results to the setup of \\cite{fernandez09a} and compare our estimates with the results of their simulations. Finally, our results are discussed in Sect.~7 and summarized in Sect.~8. ",
        "conclusions": "\\subsection{SASI or neutrino-driven convection? }  \\cite{fernandez09b} argued that the large amplitude $l=1$ oscillations appearing in their numerical simulations including iron dissociation and a heating function is due to neutrino-driven convection rather than SASI, since SASI is stabilized by iron dissociation according to \\cite{fernandez09a}. Does iron dissociation at the shock really prevent SASI from growing to large amplitudes in a realistic core collapse? In realistic simulations the compression factor at the shock can reach $\\sim 10 $, which corresponds to $ \\epsilon = 0.14 $ in the present study and a SASI amplitude of $6\\%$ of the shock distance $( r_{\\rm sh} - r_{\\rm *} ) $, quite smaller than without dissociation ($40\\%$). We point out however that a significant fraction of this amplitude decrease may be an artifact of the parameterization,  which changed the cooling function by a factor 127 in order to keep the ratio $ r_{\\rm sh}/r_{\\rm *} $ constant. This parameterization has the great advantage of being insensitive to geometrical effects that may arise if the aspect ratio between the shock and the cooling surface is changed. However cooling is then artificially low when dissociation is taken into account without heating. As a consequence, the resulting flow profile may not be more realistic than the flow profile without dissociation. As our analysis suggests that entropy gradients play an important role in the saturation of SASI, we investigated the effect of keeping the cooling function constant when dissociation is varied, resulting in a change of the shock radius (from $2.5r_{*}$ to $1.46r_{*}$, for $0<\\epsilon < 0.2v_{\\rm ff}^{2}$). By performing the same analysis as in Section~6, we then find that the saturation amplitude of SASI should decrease significantly less than when the shock radius is kept constant : $\\Delta r/( r_{\\rm sh} - r_{\\rm *} ) $ decreases by a factor $1.75$ only. Equivalently, $\\Delta r/r_{\\rm sh} $ and $\\Delta r/r_{\\rm *} $ are decreased by a factor $3.5$ and a factor $6$ respectively (to be compared with a decrease of 15 when $r_{\\rm sh}$ is constant). Although the geometrical effects make a direct comparison difficult, the fact that all these numbers are significantly smaller than the former variation by a factor 15 suggests that the decrease of the cooling function, necessary to keep the shock radius constant, plays a key role in decreasing the saturation amplitude. More insight on this issue may be gained by including the effect of neutrino heating in a parameterized manner, such that dissociation could be varied while both the cooling function and the shock radius are constant. This calculation is left for a future study.  The numerical simulations by  \\cite{scheck08} suggest that SASI is able to grow to large amplitudes even in the presence of dissociation. These simulations are significantly more realistic than the set up studied here, since they include a realistic equation of state where dissociation is taken into account in a physical way, and a simplified treatment of neutrino heating and cooling. They also differed from those by \\cite{fernandez09a} by their choice of a moving inner boundary mimicking the proto neutron star contraction. In some of the models of this article (e.g. W00), neutrino-driven convection was artificially suppressed but still SASI oscillations could grow to non negligible amplitudes. %, contradicting the claim that dissociation prevents SASI from growing to large amplitudes. It is however difficult to determine which difference between the two setups affects most importantly the saturation amplitude of SASI.  Incidentally, it is worth noting that RTi mushrooms have been identified growing on the SASI entropy gradients in Fig.~7 of \\cite{scheck08} and were interpreted in their Sect.~6.1 as secondary convection. Although in that article convection was not recognized as an agent of SASI saturation, the fact that the RTi appears at a SASI amplitude close to the saturation amplitude is consistent with a parasitic mechanism of saturation.  The interaction between SASI and neutrino-driven convection is complex and still poorly understood. Could neutrino-driven convection prevent the growth of SASI by breaking its mode structure ? One may argue that neutrino-driven convection does not feed upon the SASI mode energy, but rather converts free energy from the stationary gradients into vorticity. Could neutrino-driven convection feed SASI, either by creating vorticity which would enter the advective-acoustic cycle, or by creating sound waves \\citep{fernandez09b}? These difficult questions are beyond the scope of our study.   \\subsection{Distinguishing RTi from KHi in the simulations}  The RTi is often characterized by finger-like or mushroom-like structures as in Fig.~\\ref{imageRT}, while the KHi is characterized by vortices as in Fig.~\\ref{imageKH} and \\ref{toyfigure_KH}. However, in a complex flow containing both entropy gradients, shear and advection, RTi mushrooms may look like vortices (Fig.~\\ref{toyfigure_RT}).  The following criterion may be more useful: the RTi should occur preferentially where the entropy perturbation is maximum, while the KHi occurs where the shear is maximum. In a sloshing mode these two maxima are very distinct: the entropy oscillation is maximum at the pole (where the shock speed is maximum), while vorticity is maximum at the equator (where the inclination of the shock is maximum). The parasitic structures visible in the simulations by \\cite{scheck08} and in the movies published online by \\cite{fernandez09a} are more vigorous near the pole, in agreement with our analysis.  Furthermore, the RTi structure should grow preferentially on the half wavelength of a SASI mode where the entropy gradient is negative. By contrast, the KHi should grow on the whole extent of the SASI wavelength. An inspection of Fig.~7 of \\cite{scheck08} and of the movies by \\cite{fernandez09a} confirms this distinct feature of the RTi.   \\subsection{Numerical resolution needed to resolve the parasites}  An interesting concern raised by this saturation mechanism is that simulations should be able to resolve the parasitic instabilities properly in order to give reliable results on the nonlinear behavior of SASI. The RTi is a short wavelength instability, but as is shown in Sect.~4.3 advection tends to stabilize the small scales and makes the RTi dominated by large scales. Entropy stratification on the other hand favors small scales. As in the set up studied here the RTi is found to develop where both stabilizing effects are important, it is hard to make any prediction for its dominant wavelength. Any convergence study should verify that the grid size allows for the growth of parasites.  As an example, \\cite{scheck08} witnessed the growth of Rayleigh-Taylor mushrooms with a typical angular scale of $ l \\sim 20-30 $. They were able to capture this small angular scale by using $360$ angular zones for $180\\,^{\\circ} $. Most 2D simulations use a resolution with $>100-200$ angular zones, and would probably resolve these scales (200 zones in \\cite{murphy08b}, 121 in \\cite{burrows06}, 128-192 in \\cite{marek09}, and 60-120 zones in \\cite{ohnishi06}). However 3D simulations may not be able to resolve such small scales. For example \\cite{iwakami08, iwakami09} mostly use a resolution of $30$ angular zones for $180\\,^{\\circ}$, which may be too coarse to capture such a small scale behavior. We note that the saturation amplitude of the low-$l$ modes ($l=1-3$) in  Fig.~16 of \\cite{iwakami08} is slightly smaller at ``high resolution\" (60 zones) than at ``low resolution\" (30 zones).  While this may be explained by a suppression of parasitic instabilities at low resolution, we cannot exclude that a different saturation process may take place in 3D, as discussed below.   \\subsection{Effects of other physical ingredients}  The saturation mechanism described in this paper can be used to anticipate the effects of many other physical ingredients of the core collapse model (e.g. 3D versus 2D, the rotation rate, the magnetic field) although a detailed analysis is beyond the scope of this paper.  \\begin{itemize} \\item \\emph{3D versus 2D}: 3D simulations allow for non-axisymmetric modes that are artificially forbidden in axisymmetric simulations, thus a greater number of modes is available to the linear development of SASI. A single mode $l=1$, $m=0$ often dominates in 2D, whereas $3$ modes $l=1$, $m=0,\\pm1$ have the same growth rate in 3D if the collapsing core does not rotate \\citep{foglizzo07}. Besides, \\cite{iwakami08} found that the saturated mode amplitude is independent of $m$. We cannot exclude that nonlinear processes associated with the coupling between different mode, ignored in our analysis, are more important in 3D than in 2D. Contrary to \\cite{iwakami08} however, \\cite{blondin07a} found that \\emph{one} spiral mode dominates the 3D dynamics.  Assuming the parasitic growth of instabilities is the dominant saturation mechanism, our analysis based on a linear description of the parasites would predict the same saturation amplitudes of SASI in 2D or 3D. However the nonlinear behavior of the RTi is known to differ in 3D and 2D (e.g. \\cite{goncharov02}, \\cite{cabot06}), and this may affect the saturation of SASI. \\cite{iwakami08} reported a  smaller saturation amplitude of each individual SASI mode in 3D as compared to 2D, although the numerical convergence of this result should be further checked (Section 7.3). If confirmed, it would raise the following questions : is this difference in amplitude a consequence of the different non linear Rayleigh-Taylor behavior in 3D? Or is this the signature of  a different saturation mechanism based on the interaction of $m\\ne 0$ modes ? A more systematic parametric study, similar to \\cite{fernandez09a} but in 3D, could help check the relevance of parasitic instabilities in 3D.  \\item \\emph{Rotation rate}: \\cite{yamasaki08} have shown that rotation increases the growth rate of the spiral modes rotating in the same direction as the steady flow, while stabilizing the counter-rotating ones. If the rotation is strong enough, a single spiral mode dominates the evolution of SASI \\citep{blondin07a, iwakami08}. According to our analysis (Eq.~(\\ref{satRT})), the larger growth rate of the spiral mode could lead to a larger saturation amplitude of SASI. Nevertheless, a detailed calculation using the exact entropy and vorticity profiles of the SASI eigenmodes in a rotating flow is required in order to make an accurate prediction.  \\item \\emph{Magnetic field strength}: The effect of the magnetic field on the linear phase of SASI is yet to be understood \\citep{guilet10}, but its effect on parasitic instabilities can already be anticipated from the point of view of the magnetic tension which tends to prevent motions that distort the magnetic field lines. This effect is stabilizing for the perturbations with a wave vector parallel to the magnetic field, but does not affect those whose wave vector is perpendicular.  One would then expect that  the magnetic field does not change the maximum RTi growth rate, but selects RTi modes with a wave vector perpendicular to the field lines. In contrast, the KHi can be suppressed if the magnetic field along the direction of the transverse velocity is strong enough. In a situation where the KHi were the dominant parasitic instability, a magnetic field could potentially allow for a larger saturation amplitude. \\end{itemize}"
    },
    "0910/0910.0875_arXiv.txt": {
        "abstract": "For more than 140 years the chemical composition of our Sun has been considered typical of solar-type stars. Our highly differential elemental abundance analysis of unprecedented accuracy ($\\sim0.01$ dex) of the Sun relative to solar twins, shows that the Sun has a peculiar chemical composition with a $\\approx 20$\\% depletion of refractory elements relative to the volatile elements in comparison with solar twins. The abundance differences correlate strongly with the condensation temperatures of the elements. A similar study of solar analogs from planet surveys shows that this peculiarity also holds in comparisons with solar analogs known to have close-in giant planets while the majority of solar analogs without detected giant planets show the solar abundance pattern. The peculiarities in the solar chemical composition can be explained as signatures of the formation of terrestrial planets like our own Earth. ",
        "introduction": "For many years people have wondered about how our Sun compares to other stars and to whether our existence is related to special properties of our solar system. Angelo Secchi compared the Sun to many bright stars (Secchi 1868). He classified the stars in three types: type I which is the modern class A and early F, type II which are M stars, and type III comprising modern class G, K and early F, and called by Secchi {\\it tipo solare} (solar type), due to their spectroscopic similarity to the Sun. Furthermore, he concluded that ``{\\it le stelle di questo terzo tipo mostrano di avere una composizione identica a quella del nostro Sole}'' (Secchi 1868), meaning that solar type stars have identical composition to our Sun. Further works (e.g. Payne 1925; Edvardsson et al. 1993; Reddy et al. 2003) have not conclusively shown whether the Sun has a normal composition or not, due to the relatively large ($\\gtrsim$ 0.05) remaining systematic errors (Gustafsson 2008). Thus, for more than 140 years father Secchi's conclusion on the ``universal'' solar composition of the Sun has remained valid. In order to make further progress, it is important to eliminate many of the systematic errors ($\\sim$ 0.05-0.1 dex) that plague stellar chemical composition analyses (Asplund 2005). Solar analogs, which are G0-G5 dwarfs, and solar twins, stars almost identical to the Sun (Cayrel de Strobel 1996), are important in this context, in particular solar twins, because due to their similarity to the Sun a highly differential analysis will cancel most systematic errors.  Thus, the first step in accurate comparisons of the Sun to other stars is to find solar twins. After many years of intensive search (see review by Cayrel de Strobel 1996), Porto de Mello \\& da Silva (1997) found the first solar twin, 18 Sco [HD 146233], which is much closer to the Sun than previous candidates like 16 Cyg B [HD 186427] (see e.g. Fig. 2 of Mel\\'endez et al. 2006). More recent surveys have largely increased the number of solar twins in the field (Mel\\'endez et al. 2006; Mel\\'endez \\& Ram\\'{i}rez 2007; Takeda et al. 2007; Mel\\'endez et al. 2009; Ram\\'{i}rez et al. 2009) as well as in the open cluster M67 (Pasquini et al. 2008). The most productive survey at high resolution (R $\\sim$ 60,000 - 110,000) is being undertaken by our group. Most Northern solar twin candidates were observed with the 2.7m telescope at the McDonald observatory, and complemented with data obtained at the Keck observatory. The Southern targets were observed using the Magellan Clay 6.5 m telescope at Las Campanas, complemented with recent (August 2009) VLT observations.  Our solar twin project started in 2002, when only one solar twin was known. In order to improve our chances of finding solar twins, we first expanded the color-temperature relations of Alonso et al. (1996) to other photometric systems (Mel\\'endez \\& Ram\\'{i}rez 2003), allowing us to use existing photometry in other systems (e.g. Geneva) to select the best targets. We later improved the calibrations adding more stars and including new homogeneous systems like Tycho ($B_T - V_T$) and 2MASS (Ram\\'{i}rez \\& Mel\\'endez 2005). Although our first solar twin proposal in 2004 did not fly, the same year a more ``exciting'' proposal on Li in halo stars was granted a few nights with Keck. A small amount of that observing time was devoted to twin candidates, resulting in the discovery of the second solar twin (HD 98618; Mel\\'endez et al. 2006) and revealing that our temperature scale (and that of Alonso et al. 1996), although precise, have probably a zero-point issue. Our new selection of candidates from the Hipparcos catalog took into account a preliminary zero-point offset, which has been recently confirmed by analyses of solar twins and used for improved temperature calibrations (Casagrande et al. 2009).  New solar twin proposals at McDonald and Magellan at Las Campanas (through Australian access) during 2006 were successful. The first observing run at McDonald in April 2007 allowed us to identify the best solar twin known to date, HIP 56948 (Mel\\'endez \\& Ram\\'{i}rez 2007), which is not only very similar to the Sun physically, but has also a low Li abundance similar to solar. Its status of best solar twin has been recently confirmed by Takeda \\&  Tajitsu (2009). On the other hand, the Magellan observations at Las Campanas are opening new windows for astrophysics of the 0.01 dex level in chemical abundance accuracy. New observations taken recently at the VLT with UVES and CRIRES, promise to achieve even better precision (0.005 dex, $\\sim$1\\%), and to use solar twins to set tight constraints on Li and Be depletion in the Sun (e.g. do Nascimento et al. 2009).   \\begin{figure} \\begin{center} \\includegraphics[width=3.4in]{f01.eps} \\caption{Differences between [X/Fe] of the Sun and the mean values in the solar twins as a function of $T_{\\rm cond}$. The abundance pattern shows a break at $T_{\\rm cond}$ $\\sim 1200-1250$ K. The solid lines are fits to the abundance pattern, while the dashed lines represent the standard deviation from the fits. The low element-to-element scatter from the fits for the refractory ($\\sigma = 0.007$ dex) and volatile ($\\sigma = 0.011$ dex) elements confirms the high precision of our work. Observational errors (including errors in both the Sun and solar twins) are shown with dotted error bars, while the errors in the mean abundance of the solar twins are shown with solid error bars. } \\label{fig1} \\end{center} \\end{figure} ",
        "conclusions": ""
    },
    "0910/0910.3667_arXiv.txt": {
        "abstract": "We compute the cross-correlation between a sample of 14,000 radio-loud AGN (RLAGN) with redshifts between 0.4 and 0.8 selected from the Sloan Digital Sky Survey and a reference  sample  of 1.2 million luminous red galaxies in the same redshift range. We quantify how the clustering of radio-loud AGN depends on host galaxy mass and on radio luminosity. Radio-loud AGN are clustered more strongly on all scales than control samples of radio-quiet galaxies with the same stellar masses and redshifts, but the  differences are  largest on scales less than $\\sim 1$~Mpc. In addition, the clustering amplitude of the RLAGN varies significantly with radio luminosity on scales less than $\\sim 1$~Mpc. This proves that the gaseous environment of a galaxy on the scale of its dark matter halo, plays a key role in determining not only the probability that a galaxy is radio-loud AGN, but also the total luminosity of the radio jet. Next, we compare the clustering of radio galaxies with that of radio-loud quasars in the same redshift range. Unified models predict that both types of active nuclei should cluster in the same way. Our data show that most RLAGN are clustered more strongly than radio-loud QSOs, even when the AGN and QSO samples are matched in both black hole mass and radio luminosity. Only the most extreme RLAGN and RLQSOs in our sample, with radio luminosities in excess of $\\sim10^{26}$~W~Hz$^{-1}$, have similar clustering properties. The majority of the strongly evolving RLAGN population at redshifts $\\sim 0.5$ are found in different environments to the quasars, and hence must be triggered by a different physical mechanism. ",
        "introduction": "In recent years, galaxy formation models have become increasingly interested in the radio AGN phenomenon, because it is hypothesized that these objects may regulate the star formation history and mass assembly of the most massive galaxies and black holes in the Universe. Nearby radio galaxies in clusters are observed to inject a significant amount of energy into the surrounding gas. As the radio jets expand and interact with the surrounding medium, they are believed to heat the gas and prevent further accretion onto the central galaxy.  The precise conditions that determine whether an AGN develops radio jets/lobes are still a matter of debate. Several studies have shown that the probability for a galaxy to become radio-loud is a strong function of stellar mass and redshift (e.g. \\citealt{best05}; \\citealt{donoso}). The role that the environment plays in triggering or regulating the RLAGN phenomenon is not as well established.  \\citet{ledlow} found that the fraction of radio sources and the shape of the bivariate radio-optical was the same for objects in cluster and field environments. \\citet{best07} found that group and cluster galaxies had similar radio properties to field galaxies, but the brightest galaxies at the centers of the groups where more likely to host radio-loud AGN than other galaxies of the same stellar mass. In the local universe, \\citet{mandelbaum} analyzed a large sample of RLAGN at z$\\sim$0.1. They showed that RLAGN inhabit massive dark matter halos ($>10^{12.5}$~M$_{\\odot}$) and that, at fixed stellar mass, radio-loud AGN are found in more massive dark matter halos than control galaxies of the same mass selected without regard to AGN properties. This result implies that RLAGN follow a different halo mass - stellar mass relation than normal galaxies. \\citet{mandelbaum} also found that the boost towards larger halo masses did not depend on radio luminosity. \\citet{hickox} investigated the clustering in a small sample of higher redshift radio-loud AGN selected from the AGN and Galaxy Evolution Survey (AGES). They found no difference in the clustering amplitude of radio galaxies when compared to normal galaxies matched in redshift, luminosity and color.  Most nearby RLAGN lack any of the standard accretion-related signatures that would indicate that their black holes are growing significantly at the present day (\\citealt{hardcastle06}). In contrast, quasars are thought to be powered by supermassive black holes accreting at close to the Eddington rate. Large redshift surveys like the Two Degree Field Galaxy Redshift Survey (2dFGRS) and the Sloan Digital Sky Survey (SDSS) now provide angular positions, accurate photometry and spectra for tens of thousands of QSOs. Recent determinations of the quasar two-point correlation function have demonstrated that at $z<2.5$ quasars cluster like normal $L_*$ galaxies (\\citealt{croom}; \\citealt{coil}) and populate dark matter halos of $\\sim 10^{12}$~M$_{\\odot}$, with the clustering only weakly dependent on luminosity, color and virial black hole mass (\\citealt{shen1}).  As one moves out in redshift, the radio-loud AGN population evolves very rapidly in radio luminosity. Whether the RLAGN population also evolves strongly in black hole accretion rate, is considerably less clear. In particular, our understanding of whether there is a relationship between powerful, high redshift radio-loud AGN and quasars is quite sketchy. Around 10\\% of the quasar population is radio-loud. Numerous investigations have found that radio-loud quasars and at least {\\em some} powerful radio galaxies share a number of common characteristics, such as excess infrared emission, comparable radio morphologies and luminosities, optical emission lines, large evolutionary rates, and host galaxies with similar properties. It has thus been tempting to link both phenomena under the hypothesis that they are the same active nuclei viewed at different orientations (e.g. \\citealt{barthel89}; \\citealt{urry}).  A few facts are believed to be key in any attempt to understand the transition from the population of low-luminosity radio AGN produced by weakly accreting black holes at low redshifts, to a population of high-luminosity radio AGN that may be produced by strongly accreting black holes at high redshifts. \\citet{fanaroff} found an important correlation between radio morphology and radio power: low luminosity sources (Fanaroff-Riley Class I, FRI) show emission peaking close to the nuclei that fades toward the edges, whereas more luminous sources (Fanaroff-Riley Class II, FRII) are brightest toward the edges. \\citet{hine} discovered that radio galaxies could also be classified according to the strength of their optical emission lines: low excitation (weak-lined) radio galaxies or LERGs, and high excitation (strong-lined) objects or HERGs. Modern unification models usually associate quasars with the most powerful HERGs, and low luminosity LERGs with BL Lac objects. Although there is a notable correspondence between RLAGN luminosity, morphology and spectral type, i.e. lower luminosity FRIs with LERGs, and higher luminosity  FRIIs with HERGs, the correlations between these properties are not straightforward. There are populations of FRI sources with high excitation nuclear lines, and conversely, FRII galaxies with low excitation spectra are also common.  It has been known for years that very high redshift ($z>2$), powerful radio galaxies are often surrounded by galaxy overdensities with sizes of a few Mpc (e.g. \\citealt{pentericci}; \\citealt{miley}). Since we know that quasars at the same redshift are clustered like normal $L_*$ galaxies, this would seem to throw some doubt on a simple unified scheme for explaining both phenomena.  In view of this highly complex situation, a more statistical approach to comparing the properties of quasars and radio galaxies may yield further insight. In this paper we present measurements of the projected cross-correlation between a  sample of 14,000 radio-loud AGN with a median redshift of $z=0.55$ with the surrounding population of massive galaxies ($M_*>10^{11}$~M$_{\\odot}$). The large size of our  samples allows us to investigate in detail how clustering depends on stellar mass and on radio luminosity. By comparing the RLAGN clustering with results from control samples matched in redshift, luminosity and mass, we isolate the effect that the radio AGN phenomenon has on the clustering signal. We cross-correlate radio quasars drawn from the SDSS with the same reference sample of massive galaxies. Again, by using control samples matched in black hole mass and radio luminosity, we ensure that we compare RLAGN and RLQSOs in as uniform a way as possible.  This paper is organized as follows. In Section 2 we describe the surveys and samples used in this work. In Section 3 we explain the methodology adopted to calculate the two-point correlation function. Section 4 presents the results on radio-loud AGN and quasar clustering. Finally, in Section 5 we summarize our results and discuss the implications of this work.  Throughout the paper we assume a flat $\\Lambda$CDM cosmology, with $\\Omega_{m}=0.3$ and $\\Omega_{\\Lambda}=0.7$. Unless otherwise stated, we adopt $h=H_{0}$/(100 km s$^{-1}$) and present the results in units of Mpc~$h^{-1}$ with $h=1$. ",
        "conclusions": "In this work, we have successfully applied cross-correlation techniques to characterize the  environments of $\\sim 14,000$ radio-loud AGN with P$_{\\rm1.4 GHz}>10^{24}$~W~Hz$^{-1}$, selected from $\\sim 1.2$ million LRG at $0.4<z<0.8$. We have also compared the clustering of RLAGN with that of radio-loud quasars over the same redshift interval. By using control samples of radio-quiet objects matched in redshift, stellar mass and optical luminosity (or radio luminosity, when appropriate) we have isolated the effect such parameters have in influencing the clustering signal. The main results of this paper can be summarized as follows:  \\begin{itemize} \\item Radio AGN at $0.4<z<0.8$ are substantially more clustered than their parent luminous red galaxy population. Radio-loud AGN are also more strongly clustered than radio-quiet galaxies of the same stellar mass and redshift. The clustering differences are largest on scales less than 1~Mpc~$h^{-1}$. \\item Radio-loud AGN hosted by more massive galaxies are more strongly clustered than those hosted by less massive galaxies. However, the clustering {\\em difference} between RLAGN and control samples of radio-quiet galaxies is most pronounced for RLAGN in low mass hosts. \\item We study the dependence of the clustering amplitude on the luminosity of the radio source. For $r_p>1$~Mpc~$h^{-1}$ there is a weak, but significant anti-correlation with radio power. For $r_p<1$~Mpc~$h^{-1}$ the dependence of clustering amplitude on luminosity is more complex: the cross-correlation amplitude increases with luminosity up to $\\sim10^{25.3}$~W~Hz$^{-1}$, and then decreases for the most luminous radio sources in our sample. \\item We have compared the environments of radio-loud AGN and radio-loud QSOs. RLAGN are clustered more strongly than RLQSOs on all scales, indicate that they populate dark matter halos of different mass. These results hold even when the RLAGN and RLQSO samples are matched in radio luminosity and black hole mass. \\item There are indications that the very most luminous RLAGN and RLQSOs in our sample (P$>10^{26}$~W~Hz$^{-1}$) do have similar clustering amplitudes. Only at these very high radio powers are the space-densities of radio-loud quasars and radio galaxies similar. This implies that unification of the two AGN populations can only be valid above P$\\sim 10^{26}$~W~Hz$^{-1}$. \\end{itemize}  One major limitation of this study with regard to constraining AGN unification scenarios, is that it is based purely on photometric data from the SDSS, so we are unable to split our RLAGN sample into high-excitation and low-excitation sources. It is quite possible that the presence or absence of emission lines will provide the best way to define a population of radio galaxies that are clearly unified with the quasars. In this case, we would expect to find that the high-excitation radio galaxy population would cluster in a similar way to the quasars.  In addition, we note that because the parent sample of our RLAGN catalogue consists of luminous {\\em red} galaxies, it is also likely that we completely miss some number of RLAGN with bluer colors and stronger emission lines. The analysis of the RLAGN luminosity function presented in \\citet{donoso} indicates that the missing sources cannot constitute more than $\\sim$20\\% of the total RLAGN population, so will not dominate the clustering signal of the radio AGN  population as a whole. Nevertheless the quasar analogues among the radio galaxy population may still be under-represented in our analysis.  Fortunately, upcoming large spectroscopic surveys such as BOSS will target nearly complete samples of more than a million massive galaxies at $0.4<z<0.8$ and will provide optical spectra for tens of thousands of  radio galaxies. We will then be able to quantify the fraction of RLAGN of given radio luminosity that have emission lines and how the clustering depends on emission line strength.  The most definitive result to emerge from our analysis is clear proof that the environment of a galaxy on the scale of the dark matter halo in which it resides (i.e. on scales of $\\sim$1~Mpc~$h^{-1}$ and below), does play a key role in determining not only the probability that a galaxy a is radio-loud AGN, but also  the total luminosity of the radio jet. Combining our results with those of \\citet{best05}, we conclude that both black hole mass and environment must determine the radio-loud character of an active galaxy.  Our previous work also demonstrated that strong evolution of the radio AGN population only occurs above a characteristic radio luminosity of $\\sim10^{25}$~W~Hz$^{-1}$ (\\citealt{donoso}). It is very intriguing that the results in this paper indicate that this luminosity marks the break point in clustering trends, and that the radio luminosity where denser environment ceases to have a boosting influence also is of order $10^{25}$~W~Hz$^{-1}$.  Finally, the strong evolution of the radio source population at  radio luminosities above $\\sim10^{25}$~W~Hz$^{-1}$ combined with the strong clustering of this population, must imply that the heating rate of the gas in groups and clusters of galaxies is higher at redshifts $\\sim 0.5$ than it is at the present day. We intend to quantify this in more detail in upcoming work."
    },
    "0910/0910.1348_arXiv.txt": {
        "abstract": "We present a new mass estimate for the Hercules dwarf spheroidal galaxy (dSph), based on the revised velocity dispersion obtained by \\citet{Aden2009}.  The removal of a significant foreground contamination using newly acquired Str\\\"omgren photometry has resulted in a reduced velocity dispersion. Using this new velocity dispersion of $3.72 \\pm 0.91 {\\rm \\, km\\, s^{-1}}$, we find a mass of $M_{300}=1.9^{+1.1}_{-0.8} \\times\\, 10^6 \\,M_{\\odot}$ within the central 300\\,pc, which is also the half-light radius, and a mass of $M_{433}=3.7_{-1.6}^{+2.2}\\times10^6$M$_\\odot$ within the reach of our data to 433\\,pc, significantly lower than previous estimates. We derive an overall mass-to-light ratio of $M_{433}/L=103^{+83}_{-48} \\, [M_{\\odot}/L_{\\odot}]$. Our mass estimate calls into question recent claims of a common mass scale for dSph galaxies.  Additionally, we find tentative evidence for a velocity gradient in our kinematic data of $16 \\pm 3$\\,km\\,s$^{-1}$kpc$^{-1}$, and evidence of an asymmetric extension in the light distribution at $\\sim 0.5$\\,kpc. We explore the possibility that these features are due to tidal interactions with the Milky Way.  We show that there is a self-consistent model in which Hercules has an assumed tidal radius of $r_t = 485$\\,pc, an orbital pericentre of $r_p = 18.5\\pm 5$\\,kpc, and a mass within $r_t$ of $M_{\\mathrm{tid},r_t}=5.2_{-2.7}^{+2.7} \\times 10^6$\\,M$_\\odot$. Proper motions are required to test this model. Although we cannot exclude models in which Hercules contains no dark matter, we argue that Hercules is more likely to be a dark matter dominated system which is currently experiencing some tidal disturbance of its outer parts. ",
        "introduction": "Dwarf spheroidal (dSph) galaxies are believed to play an important role in the formation and evolution of much more luminous galaxies \\citep[e.g.][]{1994PASP..106.1225G}. dSphs are characterized by their low surface brightness, low total luminosity and spheroidal shapes which are consistent with their pressure-supported stellar kinematics~\\citep[e.g.][]{Grebel2003}.  Knowledge of dSph masses is essential for comparison with cosmological simulations of galaxy formation. Good mass estimates help us to establish whether the paucity in the number of observed systems (a few tens) to the number of predicted satellite haloes (several thousands) represents a fundamental failure of our cosmological model \\citep[e.g.][]{1999ApJ...524L..19M}, or whether it is simply telling us that galaxy formation is inefficient on small scales \\citep[e.g.][]{2006MNRAS.371..885R}. Recent studies have suggested that the dSph galaxies share a common mass within a certain radius \\citep{2007ApJ...667L..53W,2008Natur.454.1096S,2009arXiv0906.0341W}. If confirmed, this would be an important clue to the processes which regulate the formation of the lowest luminosity galaxies.  All mass estimates implicitly assume that contamination of the kinematic sample by foreground stars or unbound tidal tails is negligible, and that the system is in (or close to) virial equilibrium.  Since the mass of an equilibrium stellar system is proportional to its velocity dispersion squared, an over-estimate of the velocity dispersion will result in an inflated mass estimate.  The assumption of virial equilibrium may be called into question if the system is tidally disrupting ~\\citep{Oh1995,Kroupa1997,2007MNRAS.378..353K,2008ApJ...679..346M}.  The recently discovered Hercules dSph~\\citep{Belokurov2007}, with an ellipticity of $e\\sim 0.5$ \\citep{2007ApJ...668L..43C}, is an example of a galaxy in which assumptions of equilibrium may be incorrect.  In \\citet{Aden2009} we showed that the mean velocity of the Hercules dSph is embedded in the foreground dwarf star velocity distribution. In order to weed out the foreground dwarf stars we used the Str\\\"omgren $c_1$ index. This index is able to clearly identify red-giant branch (RGB) stars redder than the horizontal branch, enabling a separation of RGB stars in the dSph galaxy and foreground dwarf stars.  By weeding out the foreground contaminants we found that the dispersion for Hercules is reduced from $7.33 \\pm 1.08 {\\rm \\, km\\, s^{-1}}$ to $3.72 \\pm 0.91 {\\rm \\, km\\, s^{-1}}$.  In this letter we explore the consequences of this finding.  In Section~2, we derive a new mass for Hercules and show that it is not consistent with a common mass scale for the dSphs. In Section~3, we investigate the possibility of a velocity gradient in the kinematic data of Hercules. In Section~4, we discuss the relative importance of tides and dark matter in Hercules.  Section~5 summarises our conclusions. ",
        "conclusions": "We have calculated the mass of the Hercules dSph using the new velocity dispersion for the system obtained by \\citet{Aden2009}. We find that the mass within the volume enclosed by our observed stars is $3.7^{+2.2}_{-1.6}\\times \\, 10^6 \\, M_{\\odot}$, leading to a mass-to-light ratio of $103^{+83}_{-48} \\, [M_{\\odot}/L_{\\odot}]$. Interestingly, the mass within 300\\,pc is significantly lower than the ``common mass scale'' found by \\citet{2008Natur.454.1096S}, suggesting that Hercules does not share the halo properties seen in other dSphs.  We found tentative evidence for a velocity gradient of $16 \\pm 3$\\,km\\,s$^{-1}$kpc$^{-1}$, and evidence of an asymmetric extension in the light distribution at $\\sim0.5$\\,kpc.  We explored the hypothesis that these features are due to tidal interactions with the Milky Way. Assuming a tidal radius of $485$\\,pc, we show that a self-consistent model requires Hercules to be on an orbit with pericentre $r_p = 18.5 \\pm 5$\\,kpc, and with a mass within $r_t$ of $M_{\\mathrm{tid},r_t}=5.2_{-2.7}^{+2.7}\\times10^6$M$_\\odot$."
    },
    "0910/0910.1381_arXiv.txt": {
        "abstract": "The Interferometric Bidimensional Spectrometer (IBIS) installed at the Dunn Solar Telescope of the NSO/SP is used to investigate the morphology and dynamics of the lower chromosphere and the virtually non-magnetic fluctosphere below. The study addresses in particular the structure of magnetic elements that extend into these layers. We choose different quiet Sun regions in and outside coronal holes. In inter-network regions with no significant magnetic flux contributions above the detection limit of IBIS, we find intensity structures with the characteristics of a shock wave pattern. The magnetic flux elements in the network are long lived and seem to resemble the spatially extended counterparts to the underlying photospheric magnetic elements. We suggest a modification to common methods to derive the line-of-sight magnetic field strength and explain some of the difficulties in deriving the magnetic field vector from observations of the fluctosphere. ",
        "introduction": "\\begin{figure*}[t] \\centering \\includegraphics[width=\\textwidth]{fig3.eps} \\caption{a) The Ca II infrared line (854.2\\,nm). Displayed as $*$ are the sampled wavelength points (30) of data set 2, the $\\Diamond$ indicate those of data set 1. The shaded area indicates the wavelength and intensity regime within which bisectors for the use with the hybrid bisector-COG were calculated (see text). b) Stokes profiles of the strong magnetic features located at [11,16]\\,arcs in Fig.~\\ref{fig:mag}~c) (black solid: Stokes V/I; dashed: Stokes Q/I; dotted: Stokes U/I).} \\label{fig:wavepoints} \\end{figure*} Understanding the chromosphere and its magnetic structure is important to progress in understanding the solar atmosphere as a coupled phenomenon from the photosphere to the corona. A comprehensive physical model of the chromosphere must be capable of explaining a multitude of observations including, e.g., Calcium K grains \\citep{1968SoPh....3..367B}, the magnetic canopy \\citep{1983SoPh...87...37J}, and chromospheric oscillations \\citep{1993A&A...274..584K,1993ApJ...414..345L}. Constructing such a model is complicated by many things, and detailed observations such as those presented in this work are important to constraint this effort. The few available chromospheric diagnostics are unfortunately subject to complicated formation processes, for which -- in most cases -- the simplifying assumption of local thermodynamic equilibrium cannot be made, causing severe difficulties for the quantitative comparison between models and observations. Furthermore, fluctuations on small spatial and temporal scales are very challenging for both modeling and observation. Advances in both fields has lately led to an increase of attention paid to the solar chromosphere, sparking a controversy about its ``true'' physical nature. A recent summary of this controversy has been given by \\citet{2009A&A...494..269V}.  \\begin{figure*}[t] \\centering \\includegraphics[width=\\textwidth]{fig1.eps} \\caption{Sequence of subframes of a weak-field region in the Ca II infrared line core, in comparison to the synthesized model that was degraded with the point spread function of the Dunn Solar Telescope (right panel). Intensity is scaled to the same range in all panels. } \\label{fig:corr} \\end{figure*} Recent high-resolution observations suggest that in the quiet Sun there exists a weak-field domain below the classical canopy \\citep{2006A&A...459L...9W}. The prominent magnetic fields of the canopy are major constituents of the chromosphere in the stricter sense as it is seen in the H$\\alpha$ line. The weak-field domain below is referred to as ``fluctosphere'' hereafter -- a term first used by \\citet{2008IAUS..247...66W} \\sven{and previously named ``clapotisphere'' \\citep{1991SoPh..134...15R, 1995ESASP.376a.151R}.} The fluctosphere in numerical radiation hydrodynamic models is generated by interference of (acoustic) shock waves, which are excited in the photosphere and propagate upwards into the layer above \\citep[cf.][]{1994chdy.conf...47C}.  At heights around 1000\\,km above $\\tau=1$ this results in an apparent temperature pattern that has similarity to reversed granulation, yet changes with typical timescales between 20--30\\,s \\citep{2004A&A...414.1121W}. This pattern is the result of the steep temperature increase at the vertices of the interfering wavefronts, whereas the enclosed post-shock regions are cool on average.  The recent ability to synthesize spectra from models depicting the fluctosphere \\citep{2009SSRv..144..317W} allow their direct comparison to observations of this atmospheric region, under both morphological as well as dynamical aspects. Acoustic shocks have been analyzed in detail by \\citet{2009A&A...494..269V}, who use the Ca II infrared line at 854.2\\,nm as diagnostic of what they call the ``mid-chromospheric'' region of the solar atmosphere. It has been shown by \\citet{2008A&A...480..515C} that the Ca II infrared triplet provides a convenient and accessible diagnostic of the chromosphere. As \\citet{2001ApJ...554..424J}, \\citet{2009A&A...494..269V} notice a strong influence of magnetic field on acoustic processes within that regime of the solar atmosphere, and conclude that this influence may be larger than generally expected. Nevertheless, they also confirm the existence of regions with shock behavior that is similar to what is predicted by numerical models with no or at least weak magnetic field only.  A fundamental knowledge of the chromospheric magnetic field topology and its dynamics thus seems to be the next unavoidable step towards a comprehensive understanding of the solar atmosphere. However, the detailed magnetic properties of the ``mid-chromosphere'' are widely unknown so far because of the obstacles that need to be overcome in observational techniques. This in particular applies to the weak-field regions in the quiet Sun. It has been suggested that magneto-acoustic heating plays a role in the chromospheric heating \\citep[e.g.][]{2005ApJ...631.1270H, 2008ApJ...680.1542H}, even though the predominant mechanism is still debated.  Here we present and interpret for the first time high spatially resolved spectro-polarimetric measurements in the fluctosphere above strong photospheric magnetic features. We will analyze the fine-structure of the line wing and core intensity in the Ca~II infrared line in comparison to synthetic maps from numerical models. The paper is structured as follows. In Sect.~\\ref{sec:data}, we present the data observed along with the methods used for calibration. Section~\\ref{sec:top} describes the algorithm and result of the analysis of the three-dimensional topology inherent to a magnetic structure found in our data, whereas in Sect.~\\ref{sec:seq} we show the results of our time sequence analysis. We conclude with a summary of our results in Sect.~\\ref{sec:con}. ",
        "conclusions": "\\label{sec:con} Spectro-polarimetric observations of the solar atmosphere with a focus on heights equivalent to the low chromosphere have been presented in this work. We find that fluctospheric regions of the solar atmosphere, i.e. the weak field domain between the photosphere and the magnetically dominated chromosphere, display an intensity pattern that shows structures of similar spatial scales as reversed granulation but is evolving on much faster time scales. From comparison to the predicted synthesized intensity of three-dimensional radiation hydrodynamic models of the fluctosphere, we conclude that the pattern is the intensity signature of propagating, interacting shock waves. The pattern's dynamic time scale of 59\\,s is comparable to the 20--30\\,s of temperature cuts at the height of 1000\\,km above $\\tau_{cont}=1$, when taking into account that seeing and finite spectral passband of the instrument will lead to prolonged correlation times.  We propose a method to derive the 3D topology of the absolute LOS magnetic field flux per pixel. The method works for simple cases in non-emitting lines like the Ca II IR line in the Quiet Sun. For such cases, we were capable to reconstruct the topology of a strong magnetic structure that extends from photosphere to chromosphere. The structure becomes weaker with height, but does develop a filament like structure when recovering the field using the line core. The method is not capable to infer filling factors, and thus cannot reproduce the field strength of an unresolved magnetic element.  A complementary dataset with higher signal-to-noise ratio shows structures in total linear polarization encirceling strong magnetic elements. This strongly suggests a flux tube funnel structure in the chromosphere. However, the location of the horizontal field remains unclear, and a static flux tube funnel appears to be a too simplified model for magnetic fields in the chromosphere.  Highly spatially resolved measurements of the magnetic fields in the chromosphere are important when trying to gather a deeper understanding of the chromospheric energy balance. The work presented here clearly shows that, despite the usage of advanced techniques, the signal-to-noise ratio of currently available data prohibits the detection of weak field strengths and definitely not at the temporal resolution which is implied by the intensity variations seen in chromospheric/fluctospheric diagnostics such as the Ca IR line core. Ultimately, telescopes with significantly larger apertures are needed to get a more detailed picture of the chromosphere. The difficulties of the observation of weak magnetic fields in the chromosphere are only surpassed by the interpretation of such data. In general, the height origins of the Stokes signals are unclear and do not only depend on contribution function but also on height distribution of the magnetic field strength. In addition, discontinuities and steep gradients in magnetic field strength are difficult to detect and remain a fundamental problem of all methods to recover the magnetic field vector from measurements with methods available today."
    },
    "0910/0910.1512_arXiv.txt": {
        "abstract": "Analytic distribution functions (\\df s) for the Galactic disc are discussed. The \\df s depend on action variables and their predictions for observable quantities are explored under the assumption that the motion perpendicular to the Galactic plane is adiabatically invariant during motion within the plane. A promising family of \\df s is defined that has several adjustable parameters. A standard \\df\\ is identified by adjusting these parameters to optimise fits to the stellar density in the column above the Sun, and to the velocity distribution of nearby stars and stars $\\sim1\\kpc$ above the Sun. The optimum parameters imply a radial structure for the disc which is consistent with photometric studies of the Milky Way and similar galaxies, and that 20 per cent of the disc's luminosity comes from thick disc.  The fits suggest that the value of the $V$ component of the Sun's peculiar velocity should be revised upwards from $5.2\\kms$ to $\\sim11\\kms$. It is argued that the standard \\df\\ provides a significantly more reliable way to divide solar-neighbourhood stars into members of the thin and thick discs than is currently used. The standard \\df\\ provides predictions for surveys of stars observed at any distance from the Sun.  It is anticipated that \\df s of the type discussed here will provide useful starting points for much more sophisticated chemo-dynamical models of the Milky Way. ",
        "introduction": "A major thread of current research is work directed at understanding the origin of galaxies. There are excellent prospects of achieving this goal by combining endeavours in three distinct areas: observations of galaxy formation taking place at high redshift, numerical simulations of the gravitational aggregation of dark matter and baryons, and studies of the Milky Way. The latter field is dominated by a series of major observational programs that started fifteen years ago with ESA's Hipparcos mission, which returned parallaxes and proper motions for $\\sim10^5$ stars \\citep{Perryman}. Hipparcos established a more secure astrometric reference frame, and the UCAC2 catalogue \\citep{UCAC} uses this frame to give proper motions for several million stars. These enhancements of our astrometric database have been matched by the release of major photometric catalogues [DENIS \\citep{DENIS}, 2MASS \\citep{2MASS}, SDSS \\citep{SDSS}] and the accumulation of enormous numbers of stellar spectra, starting with the Geneva--Copenhagen Survey \\citep[][hereafter GCS]{Nordstrom04,HolmbergNA} and continuing with the SDSS, SEGUE \\citep{Yanny} and RAVE \\citep{Steinmetz} surveys -- on completion the SEGUE and RAVE surveys will provide $0.25\\times10^6$ and $\\sim10^6$ low-dispersion spectra, respectively. These spectra yield good radial velocities and estimates of $\\feh$ and $\\afe$ with errors of $\\sim0.1\\dex$.  Two surveys (HERMES and APOGEE) are currently being prepared that will obtain large numbers of medium-dispersion spectra from which abundances of significant numbers of elements can be determined. The era of great Galactic surveys will culminate in ESA's Gaia mission, which is scheduled for launch in late 2011 and aims to return photometric and astrometric data for $10^9$ stars and low-dispersion spectra for $>10^7$ stars.  The Galaxy is an inherently complex object, and the task of interpreting observations is made yet more difficult by our location within it. Consequently, the ambitious goals that the community has set itself, of mapping the Galaxy's dark-matter content and unravelling how it was assembled, can probably only be attained by mapping observational data onto sophisticated models. We are developing a modelling strategy that has as its point of departure analytic approximations to the distribution functions (\\df s) of various components of the Galaxy (McMillan et al.\\ in preparation).  In this paper we present such approximations for the thin and thick discs.  The paper is organised as follows. Section 2 explains how the \\df\\ is assembled. Section 3 compares the \\df's predictions for various observables to data. In particular evidence is presented that the Sun's $V$ velocity is conventionally underestimated by $\\sim6\\kms$ and predictions are given for velocity distributions as a function of distance from the plane.  Evidence is presented that  the  standard \\df\\ provides a cleaner division of solar-neighbourhood stars into members of the thin and thick discs than has been available hitherto. Section 4 sums up and looks ahead. ",
        "conclusions": "We have explored the ability of distribution functions to provide models of the thin and thick discs of the Milky Way. Our \\df s are analytic functions of the actions of orbits, which ensures that there is an intuitive relation between the observable properties of the population a \\df\\ describes and the functional form of the \\df, and a meaningful way to compare models that use different gravitational potentials. In this paper we have used expressions for the actions that are only approximate, and imply that a star's vertical motion is adiabatically invariant during the star's motion parallel to the plane. In a forthcoming paper (McMillan et al., in preparation) orbital tori will be used to eliminate this approximation, and thus quantify its validity.  We have shown that the vertical density profile and kinematics of the disc are accurately modelled by the extremely simple \\df\\ (\\ref{powerdf}). However, we rejected this \\df\\ because  it is essential to be able to break the \\df\\ for the thin disc down at least into contributions from stars of various ages, and ideally into contributions from ranges in both age and metallicity. That is, we must recognise that the Galaxy is built up of innumerable stellar populations of various ages and metallicities, and each population has its own \\df. In this paper we have only begun to explore the resulting complexity by ascribing a single \\df\\ to the thick disc and modelling the thin disc as a superposition of \\df s for stars of different ages. In reality both  discs are chemically inhomogeneous and we should assign a distinct \\df\\ to the stars born at each time with each chemical composition \\citep[e.g.][]{SchoenrichBII}. Hence the \\df s presented in this paper should be considered building blocks  from which more elaborate \\df s may be in due course constructed.  Our most basic building block is a ``pseudo-isothermal'' population of stars. \\figref{fig:HF} shows that the vertical distributions of young stellar populations is well modelled by a pseudo-isothermal population. The density of a pseudo-isothermal population does not decline exponentially with $z$, but Figs.~\\ref{fig:thinD} and \\ref{fig:thinthick} show that, remarkably, the composite population produced by stochastic acceleration of stars does have an exponentially decreasing density profile. An excellent fit to the observed density profile of the entire disc is obtained when a pseudo-isothermal thick disc is added to the composite thin disc. The dispersion in $v_z$ of thin-disc stars increases from $17.4\\kms$ in the plane to $33\\kms$ at $2.5\\kpc$, while that of the thick-disc stars increases from $\\sim 35\\kms$ in the plane to $\\sim48\\kms$ at $2.5\\kpc$. The thick disc contributes to the solar cylinder 24 per cent of the luminosity contributed by the thin disc, or 19.4 per cent of the total luminosity of the disc.  Even though we are assuming that the dynamical coupling between motions in and perpendicular to the plane is weak, two features of our \\df s lead to strong correlations between distributions in $v_R$ and $v_z$. One feature is the fact that random velocities must increase as one moves inwards, and the other is the simultaneous increases in $\\sigma_R$ and $\\sigma_z$ that are driven by stochastic acceleration of a coeval population. Comparison of Figs.~\\ref{fig:DF} and \\ref{fig:thinthick} show that, on account of this correlation, the distribution of local stars in the $(v_R,v_\\phi)$ plane is atypical of the stellar population of the whole solar cylinder in just such a way that our composite disc \\df\\ can simultaneously provide reasonable matches to the very different shapes of the distributions of GCS stars in $v_R$ and $v_\\phi$. The widths of the model distributions in $v_R$ and $v_\\phi$ are controlled by a single parameter, $\\sigma_{r0}$. The shape of the $v_R$ distribution is predetermined by our choice of the \\df s functional form. The value of the parameter $R_\\d$ provides limited control of the shape of the $v_\\phi$ distribution and we obtain the best fit to the observed distribution when this parameter is chosen such that the disc's surface density declines roughly exponentially with scale length $2.5\\kpc$, which happens to agree with the scale length inferred from near-IR star counts by \\cite{Robin03}.  In principle the \\df\\ of the thick disc should be tightly constrained by the dependence on  $z$ of the velocity dispersions $\\sigma_R$ and $\\sigma_z$. These dependencies have recently been determined for SDSS stars by two groups. Unfortunately, their results seem to be incompatible and the reasons for the conflict are unknown.  The standard model provides an excellent fit to the seminal work of Kuijken \\& Gilmore, perhaps because the gravitational potential in which the \\df\\ is evaluated was partly fitted to that work. Some of the difficulties encountered here with fitting newer data may arise from inaccuracy of the potential used.  A worthwhile exercise would be to fit data from the GCS, RAVE and SDSS surveys to models that combined \\df s of the type presented here with and a multi-parameter gravitational potential: by simultaneously fitting the parameters in both the \\df\\ and the potential, one should be able to obtain reasonable fits to the data, providing the latter have been purged of such evident inconsistencies as those seen in \\figref{fig:powersig}. Data from more than one survey would probably have to be used since SDSS stars are too faint to constrain the thin disc tightly, although the RAVE survey, which certainly probes the thick disc effectively, may include enough nearby stars to make the Hipparcos-based GCS survey obsolete.  The model fit to the $v_\\phi$ distribution of GCS stars is far from perfect. Some of the disagreement arises because, as is well known, the Galactic bar and spiral arms give rise to features (``star streams'') in the local velocity distribution that are inconsistent with the Galaxy being axisymmetric and in a steady state, as our models assume.  Our favoured model $v_\\phi$ distribution would fit the data better if the conventional value of the solar motion $V_\\odot$ were $\\sim6\\kms$ too low. Tentative support for such an increase in the $V_\\odot$ is provided by astrometry of stellar masers \\citep{Reid09,McMillanB09}, and any increase would also tend to bring the Galaxy more into line with the Tully--Fisher relation between $\\vc$ and $M_I$ for external galaxies.  By systematically perturbing the velocities of all solar-neighbourhood stars, spiral structure might lead to the classical approach to the determination of $V_\\odot$ yielding an underestimate. Further work is required to explore this possibility, and at this stage we would merely stress that the systematic error in $V_\\odot$ is much larger than the formal errors given by DB98 and \\cite{AumerB}.  In the models, the asymmetric drifts of both the thin and thick discs increase with height. A disc's asymmetric drift is largely controlled by its parameter $R_\\d$ and in the standard model the asymmetric drift of the thin disc exceeds that of the thick disc above $1\\kpc$ because we have adopted a slightly larger value of $R_\\d$ for the thick disc than for the thin disc.  A popular strategy for assigning solar-neighbourhood stars to the thin or thick disc is to find the values taken by each disc's model \\df\\ at the star's location. The model \\df s used are perfectly ellipsoidal but we show that such \\df s provide poor approximations to the thick-disc component of the standard \\df, so a markedly cleaner separation of the two discs could be obtained by replacing the ellipsoidal \\df s by the thin- and thick-disc components of the  standard \\df.  Although the observational material relating to the Galaxy has increased enormously in recent years, we have shown that much of the available data can be successfully modelled with a simple analytical \\df. In a couple of aspects the data are in mild conflict with the \\df, but it is at least as likely that the fault lies with the data as the \\df. In the coming decade the volume and quality of the observational material available will increase dramatically. We anticipate that comparisons between each new data set and an evolving standard \\df\\ will reveal successes and failures similar to those encountered here. The successes will confirm the value of the \\df\\ as a summary of a large and inhomogeneous body of data, and the failures will lead to critical re-examination of both data and \\df. Sometimes the failure will arise from a defective calibration of the data or incorrect assumptions used in its reduction, and other times it will indicate that the \\df\\ is too simplistic.  Either way we will learn something new and interesting.  In this paper the \\df's parameters have been fitted to the data by eye and no attempt has been made to quantify uncertainties in parameter values. Clearly such uncertainties are important, and they could be most securely established by carrying the \\df's predictions closer to the raw observations than we have done.  In future work probability distributions in colour--magnitude--proper-motion space, etc., should be predicted that can be compared with the actual star counts.  Upcoming infrared surveys, such as the VHS with Vista and APOGEE, will probe the disc at remote locations. The predictions of the standard \\df\\ for those locations will be presented shortly, after orbital tori have been introduced as the means to convert between Cartesian and angle-action variables. This upgrade will make obsolete the approximation of adiabatically invariant vertical motions  used here."
    },
    "0910/0910.1724_arXiv.txt": {
        "abstract": "{} {We investigate the effect of  the   physical environment  on water and ammonia abundances   across the S140  photodissociation  region (PDR)   with an embedded outflow. } {We have used the Odin satellite to obtain strip maps of the ground-state rotational transitions of \\emph{ortho}-water and \\emph{ortho}-ammonia,   as well as CO(5\\,--\\,4) and  \\13co(5\\,--\\,4) across the    PDR, and  \\water18 in the central position. A physi-chemical inhomogeneous PDR model was used to  compute the   temperature and  abundance distributions for water, ammonia and CO. A multi-zone escape probability method then calculated the level populations and intensity distributions. These results are  compared to a homogeneous model computed with an enhanced version of the {{\\tt RADEX}} code. } {\\h2o, \\nh3 and \\13co show emission from an extended PDR with a narrow line width of $\\sim$3\\,\\kms. Like CO, the water line profile is   dominated by outflow emission, however,   mainly in the red wing. Even though CO shows strong self-absorption, no signs of self-absorption  are seen in the water line. \\water18 is not detected. The PDR model  suggests that the water emission  mainly arises from the surfaces of optically thick, high density clumps with $n(\\mathrm{H_2}$)$\\ga$10$^{6}$\\,\\cmcub~and a clump water abundance, with respect to H$_2$, of 5\\x10$^{-8}$. The mean water abundance in the PDR is    5\\x10$^{-9}$, and between $\\sim$2\\x10$^{-8}$\\,--\\,2\\x10$^{-7}$ in the outflow derived from a simple two-level approximation. The   {{\\tt RADEX}} model    points to  a somewhat higher   average PDR water abundance  of  1\\x10$^{-8}$. At low temperatures deep in the cloud the water emission is weaker,  likely due to adsorption onto   dust grains, while  ammonia is still abundant. Ammonia is also observed in  the extended clumpy PDR, likely   from   the same  high density and warm  clumps as water. The average ammonia abundance is   about the same as for water: 4\\x10$^{-9}$ and  8\\x10$^{-9}$ given by  the PDR model and {\\tt RADEX}, respectively. The differences between the models most likely arise due to uncertainties in density, beam-filling and volume filling of clumps. The similarity of water and ammonia PDR emission is also seen in the almost identical line profiles observed close to the bright rim. Around the central position,  ammonia also shows some outflow emission although weaker  than   water  in the red wing. Predictions of the \\h2o \\trans~and 1$_{1,\\,1}$\\,--\\,0$_{\\,0,\\,0}$ antenna temperatures  across the PDR are estimated with our PDR model for the forthcoming observations with the Herschel Space Observatory. } {} ",
        "introduction": " ",
        "conclusions": "\\label{section conclusions}  We have used the Odin satellite to observe water, ammonia and carbon-monoxide in the well-known molecular cloud S140. We have simultaneously observed   five-point strip maps across the  bright rim in S140  of the \\emph{ortho}-\\h2o \\trans~and  the \\emph{ortho}-\\nh3 1$_0$\\,--\\,0$_{\\,0}$ transitions, as well as \\mbox{CO(5\\,--\\,4)} and \\mbox{\\13co(5\\,--\\,4)}. The  \\nh3 transition has never previously been observed in S140. Observations of \\water18 in the central position resulted in a non-detection at a rms level of 8\\,mK. As support observation we also mapped \\mbox{\\13co(1\\,--\\,0)}    with the Onsala 20m telescope.  Like CO, the water line-profile  is dominated by emission from a NW\\,--\\,SE  outflow, however, mainly in the red wing. Strong self-absorption is seen in the optically thick CO emission, while no obvious signs are seen in the \\emph{ortho}-\\h2o, \\emph{ortho}-\\nh3 or the almost optically thin \\13co line profiles. In  addition to the outflow, our water line shows emission from a more extended NE\\,--\\,SW elongated PDR. Both these components originate approximately around our central position.   No additional emission closer to the bright rim or further into the molecular cloud is detected. Close to the bright rim the   temperature is most likely too high for a detection of our transition with   an upper state energy of 61\\,K. Instead, higher-lying transitions, observable with the Herschel Space Observatory, will have their peak intensity shifted towards  the bright rim \\citep{2005A&A.440.559.Poelman.Spaans, 2006A&A.453.615.Poelman.Spaans}.  Even closer to the bright rim at a few magnitudes of $A_\\mathrm{V}$, water is, however, photo-dissociated by the UV field.   The \\emph{ortho}-\\nh3 emission seems to emanate from the same high density clumps in the PDR and the outflow as water, but  also shows  additional emission further into the cloud where the ambient gas temperature drops    to  about 30\\,K. Compared to water  in the central position, ammonia has a weaker outflow emission in the red wing although similar in the blue outflow. Close to the bright rim, where the outflow contribution to the emission is very low,  the water and ammonia line profiles are almost identical, suggesting an origin in the same gas and velocity fields of the PDR. The \\13co line also shows a very similar line profile as \\h2o and \\nh3 in this position with a narrow line width of $\\sim$3\\,\\kms.  Abundances, with respect to  H$_2$, in the PDR  are estimated  both with an enhanced version of the homogeneous {\\tt RADEX} code and with a clumpy PDR model. This   model points to   low mean water and ammonia abundances  in the PDR of 5\\x10$^{-9}$ and 4\\x10$^{-9}$, respectively. In the high-density clumps both the average water and ammonia abundances increase to 5\\x10$^{-8}$. To match the observed PDR antenna temperatures  with Odin and SWAS, a clumpy medium is required by the model with a high molecular hydrogen density in the clumps of $\\ga$1\\x10$^6$\\,\\cmcub. The resulting {\\tt RADEX} mean abundances are twice as high: 1.0\\x10$^{-8}$  and 8\\x10$^{-8}$ for water and ammonia, respectively, using a molecular hydrogen density of 4\\x10$^{5}$\\,\\cmcub~and a kinetic temperature of 55\\,K. The differences most likely arise from the uncertainty in density, beam-filling, and volume filling of the clumps. The opacity    of the narrow PDR component of the \\trans~transition is constrained by the narrow line width, and is  estimated  by  {\\tt RADEX} to be $\\sim$7. The PDR  model also confirm  a  low water opacity with an unweighted mean opacity of 17   and a model range of $\\sim$10$^{-5}$\\,--\\,800. The mean outflow water abundance,  derived from a simple two-level approximation, is higher than in the PDR by at least one  order of magnitude, $\\sim$2\\x10$^{-8}$\\,--\\,2\\x10$^{-7}$.     Predictions of antenna temperatures for observations with   HIFI  are given by our PDR model of the \\emph{ortho}- and \\emph{para}-\\h2o \\trans~and 1$_{1,1}$\\,--\\,0$_{0,0}$ transitions for  nine positions across the bright rim, and are found to peak around 70\\,--\\,80\\arcsec~from the dissociation front in agreement with our observations."
    },
    "0910/0910.2668.txt": {
        "abstract": "The slope and zero-point of the unevolved main sequence as a function of metallicity are investigated using a homogeneous catalog of nearby field stars with absolute magnitudes defined with revised $Hipparcos$ parallaxes, {\\it Tycho-2} photometry, and precise metallicities from high-dispersion spectroscopy. $(B-V)$ - temperature relations are derived from 1746 stars between [Fe/H] $= -0.5$ and +0.6 and 372 stars within 0.05 dex of solar abundance; for T$_e$ = 5770 K, the solar color is $B-V$ = 0.652 $\\pm$ 0.002 (s.e.m.). From over 500 cool dwarfs between [Fe/H] = $-0.5$ and +0.5, $\\Delta(B-V)/\\Delta$[Fe/H] at fixed $M_V$ = 0.213 $\\pm$ 0.005, with a very weak dependence upon the adopted main sequence slope with $B-V$ at a given [Fe/H]. At Hyades metallicity this translates into $\\Delta M_V/\\Delta$[Fe/H] at fixed $B-V$ = 0.98 $\\pm$ 0.02, midway between the range of values empirically derived from smaller and/or less homogenous samples and model isochrones. From field stars of similar metallicity, the Hyades ([Fe/H] = +0.13) with no reddening has $(m-M)_0$ = 3.33 $\\pm$ 0.02 and M67, with E$(B-V)$ = 0.041, $A_V$ = 3.1E$(B-V)$, and [Fe/H] = 0.00, has $(m-M)_0$ = 9.71 $\\pm$ 0.02 (s.e.m), where the errors quoted refer to internal errors alone. At the extreme end of the age and metallicity scale, with E$(B-V) = 0.125 \\pm 0.025$ and [Fe/H] $= +0.39 \\pm 0.06$, comparison of the fiducial relation for NGC 6791 to 19 field stars with $(B-V)$ above 0.90 and [Fe/H] = +0.25 or higher, adjusted to the metallicity of NGC 6791, leads to $(m-M)_0 = 13.07 \\pm 0.09$, internal and systematic errors included. ",
        "introduction": "Distance determination to open and globular clusters is key to placing them in the proper Galactic evolutionary context and an indispensable component in evaluating stellar evolution as a function of mass, chemical composition, and age. For nearby clusters like the Hyades, Praesepe, the Pleiades and Coma, $Hipparcos$ parallaxes \\citep{esa}, coupled with proper-motion and radial-velocity memberships, have generated precise distances, independent of the cluster's composition and reddening, though these results have not been without controversy \\citep{pi98, vl99, sn01, ma02, so05, vl09}. Any remote cluster beyond the reach of parallax with a known chemical makeup comparable to a nearby cluster can be compared differentially using stars on the unevolved main sequence to obtain a reliable distance (see, e.g., \\citet{pi04, an07}). While unevolved main sequences of nearby clusters are ideal reference points due to the uniformity of age and composition among the cluster stars, the range in chemical composition sampled by the nearby clusters is approximately [Fe/H]$ = -0.2$ to +0.2. For open clusters outside this range and for all globular clusters, a traditional fallback procedure is to use cooler field dwarfs with parallaxes and well-defined abundances either to isolate a stellar sample of similar abundance or to interpolate among stars that bracket the metallicity of interest \\citep{tw99, gr03}. Given the broad gaussian distribution in [Fe/H] for field stars, this approach becomes more of a challenge for clusters whose metallicity places them in the extended, low-metallicity tail of the field star sample.  A routine alternative to field star comparison has been the construction of theoretical isochrones which can be tuned to any combination of composition and age, though ultimately these models must be linked to the empirical data of the field stars by matching the theoretical model combinations of mass, luminosity, composition and temperature to the observed values of mass, absolute magnitude, abundance and color at a given age. For stars of approximately solar mass and higher, there is reasonable agreement among the various isochrone compilations currently available in the literature and on-line, with most discrepancies among the models understandable in terms of differences in the adopted parameterizations of the internal physics, the model atmospheres, and the relations used to transfer from the theoretical to the observational plane \\citep{pad, bas04, de04, vrs06, ma08, des08}. As one might expect, the discrepancies within the observational color-magnitude diagram (CMD) plane grow larger as one moves to stars of lower mass and/or more extreme compositions, reflecting the same paucity of empirical constraints found when attempting a direct match of distant clusters to nearby field stars with parallaxes.  Issues of reddening and metallicity determination aside, the more contentious discussions of cluster distances generally have focused on metal-poor systems of the thick disk and halo, while the metal-rich end of the distribution has been moderately immune due to the rich sample of nearby stars of solar abundance, the proximity of the Hyades, and the rarity of clusters with compositions significantly higher than the Hyades. For almost 50 years the one potential exception has been the old open cluster, NGC 6791, though the extent of its anomaly has been hidden beneath disagreements over its reddening, composition, distance, and age. In recent years significant progress has been made in constraining the first two parameters, both of which are critical to defining the third using main-sequence fitting, while the fourth has no bearing on the distance if the comparison is made to stars of sufficiently low mass. In every instance where revised reddening and/or metallicity estimates have been obtained, an improved distance modulus has been derived through comparison with theoretical isochrones \\citep{bed08, ca06, ca05, ki05, stet, ch99, kr95, tr95, ga94, mj94, de92}. For the apparent distance modulus, the current range from main sequence fitting extends from a low of 13.1 \\citep{stet} to a high of 13.6 \\citep{slv03,atm07}. While the scatter is partly due to the adoption of different values for [Fe/H] and E$(B-V)$, a large portion is tied to disagreement over the location of the unevolved main sequence among the isochrones at high metallicity and the range of the CMD used to define an adequate fit to the models; in some instances, the unevolved main sequence and giant branch features are used simultaneously to optimize the fit.  To circumvent the issues presented by the discrepancies among metal-rich isochrones, it was decided that the distance to NGC 6791 could best be obtained by an empirical fit to unevolved field stars of similar composition. The feasibility of this option has been enhanced by the more restricted range among recent determinations of the cluster parameters, the availability of revised parallaxes and broad-band photometry from {\\it Hipparcos/Tycho-2} \\citep{esa,vl07}, and the compilation of a catalog of precise spectroscopic abundances, temperatures, and surface gravities on a common scale for almost 2100 nearby field stars. While the initial motivation for this study was the distance to NGC 6791, the need to define the precise location of the unevolved main sequence at high metallicity using the very restricted field star sample at comparable [Fe/H] generated a more comprehensive investigation of the metallicity dependence of the unevolved main sequence for stars of typical disk metallicity ([Fe/H] = -0.5 to +0.5). This paper builds upon the approach laid out in \\citet{pe03}, using a sample expanded by an order of magnitude to derive the change in $M_V$ for cool dwarf stars at a given $B-V$ as [Fe/H] is varied from -0.5 to +0.5. While it is usually assumed that the ratio, $\\Delta M_V/\\Delta$[Fe/H], is constant with metallicity for unevolved disk dwarfs, derived values from observation and theoretical models vary by a more than a factor of two \\citep{ko02, pe03, pi04, ka06}. The goal of this investigation is to explore the assumption of a constant ratio, derive it, and test the consistency of the field star main sequence when applied to well-studied open clusters. Unless noted otherwise, errors quoted are standard deviations.  The layout of the paper is as follows. In Sec. 2 we will describe the database of astrometry, photometry, and spectroscopy for the nearby stars that is used in Sec. 3 to define the characteristics of the unevolved main sequence as a function of color and metallicity. In Sec. 4, the derived CMD relation is used to estimate distances to the well-studied open clusters, M67 and the Hyades, as well as the extremely metal-rich cluster, NGC 6791. Sec. 5 contains a summary of our conclusions. ",
        "conclusions": "Distance determination for field stars and clusters remains a primary observational objective for those interested in understanding stellar and Galactic evolution. Ideally, parallaxes for a large sample of nearby stars with well-defined abundances and temperatures/colors would allow one to map the impact on the absolute magnitude of varying the abundance of a star of a given temperature/color. With this information in hand, determining the distance to any cluster with reliable abundance and reddening information becomes a straightforward task. Unfortunately, the reality is somewhat different. While reliable parallaxes are available for a large sample of field stars and most have colors on the {\\it Tycho-2} system, the critical component missing from the picture has been a comparable catalog of precise metallicity estimates. The observational approaches to defining the ratio of $\\Delta M_V$ with [Fe/H] \\citep{ko02, pe03, ka06} have relied upon photometric abundances coupled to a spectroscopic subset of the sample. The slope of the relation is either assumed to be constant with [Fe/H] and color and/or tested through the use of theoretical isochrones. For \\citet{ko02} and \\citet{ka06}, the baseline used to define the slope is extended to [Fe/H] below $-1$ through the inclusion of halo dwarfs and/or globular clusters. The dataset of \\citet{ka06} extends to stars where evolution off the main sequence is significant and may produce a selection bias in the mean absolute magnitude with [Fe/H].  Starting with a dataset of almost 2000 stars with reliable parallaxes, spectroscopic abundances, and homogeneous colors, we have reduced the sample to approximately 500 stars between [Fe/H] = $-0.5$ and +0.5 with colors red enough that evolution off the main sequence should be negligible. However, as evidenced by the old, metal-rich turnoff stars in NGC 6791, even our metallicity-dependent cutoff allows some significantly evolved stars into the mix. Over the primary color range of interest, $B-V$ = 0.75 to 1.15, the ratio with metallicity, $\\Delta M_V/\\Delta$[Fe/H]$ = 0.98 \\pm$ 0.02 with no evidence for a color or [Fe/H] dependence. Because this value assumes a universal slope of 4.53 for the main sequence with $B-V$, it is probably better to define the relation in terms of $\\Delta(B-V)/\\Delta$[Fe/H]$ = 0.213 \\pm$ 0.005, which is only weakly dependent upon the adopted slope of the main sequence. The concern that the slope of the main sequence with $B-V$ varies with [Fe/H] is real and not dependent upon the multiple sets of theoretical isochrones which are inconsistent on this point. From the observed cluster sequences analyzed in this investigation, the Hyades has a main sequence slope of 4.5; over the same color range, the fiducial relation for NGC 6791 has a slope of 5.4.  With the field-star relations in hand, reliable estimation of cluster distances should follow. The challenge in this phase is ensuring that the clusters are on the same photometric color and spectroscopic abundance system as the field stars which, for more distant objects, also includes accurate reddening estimation. Because our catalog includes Hyades stars and the Hyades is assumed to be reddening-free, the only potential source of controversy is the photometric scale. Our analysis confirms what has been found by others \\citep{jon06, an07, jon08}. The \\citet{joh55} photometric system in the Hyades is offset from the SAAO/E-region standards that define the transformed {\\it Tycho-2} system by approximately -0.009 mag and 0.02 mag in $B-V$ and $V$, respectively. When these corrections are applied to the fiducial relations for the Hyades \\citep{pi04}, the cluster produces an excellent match to the field stars at the appropriate [Fe/H] with the cluster modulus set at $(m-M)_0$ = 3.33.  The second test of the system is a match to M67. Because this cluster and NGC 6791 are two steps further removed from the field star sample in that there are no cluster stars with spectroscopic abundances in our catalog and the photometric data are not directly comparable to the {\\it Tycho-2} system, we have applied the same color offset found for the Hyades to the cluster data and made differential comparisons based upon a commonly adopted reddening and metallicity estimate. If the M67 data of \\citet{sa04} are matched to the Hyades fiducial relation shifted to [Fe/H] = 0.00, the apparent modulus of the cluster becomes 9.81 $\\pm$ 0.02; if a match is made directly to the field stars of identical [Fe/H], the modulus increases by 0.03. The comparable value from \\citet{an07} and \\citet{pa08} is 9.765. Note that part of the offset relative to \\citet{an07} may be a product of the elegant but hybrid approach that defines the main sequence relations using a mix of theoretical isochrones with color corrections defined by cluster data. The isochrones used by \\citet{an07} define a ratio, $\\Delta M_V/\\Delta$[Fe/H], over the color range of interest that is typically 1.4, significantly larger than our derived value near 1. With the Hyades CMD position fixed via parallax, differential shifts with metallicity then define the predicted position of clusters like M67. Even for a change in [Fe/H] of 0.13, an overestimate of 25$\\%$ in the slope would lead to an underestimate of the distance by 0.03 mag.  Finally, at the extreme end of the age and metallicity scales among open clusters, the distance to NGC 6791 is derived using field stars of comparable metallicity. Initial comparisons using the adjusted Hyades relation exhibited evidence for significant evolutionary effects off the main sequence for the bluer stars in our sample, implying that the average star in the comparison is older than the Hyades.  This trend virtually disappears when the data are compared to the fiducial relation for the cluster. In fact, the bluest stars in the field are systematically fainter than the stars in the cluster at the same color, indicating that they are, as expected, younger than NGC 6791. Adopting [Fe/H] = 0.40 and E$(B-V)$ = 0.15, we find $(m-M)_0$ = 13.13 $\\pm$ 0.09. By comparison, the definitive value tied to analysis of the eclipsing binary, V20, using the same parameters is $(m-M)_0$ = 13.00 $\\pm$ 0.10. The significance of the difference is marginal, especially given the underlying question of the photometric zero-points in $V$ and $B-V$."
    },
    "0910/0910.0980_arXiv.txt": {
        "abstract": "We consider a gravitational theory of a scalar field $\\phi$ with nonminimal derivative coupling to curvature. The coupling terms have the form $\\kappa_1 R\\phi_{,\\mu}\\phi^{,\\mu}$ and $\\kappa_2 R_{\\mu\\nu}\\phi^{,\\mu}\\phi^{,\\nu}$ where $\\kappa_1$ and $\\kappa_2$ are coupling parameters with dimensions of length-squared. In general, field equations of the theory contain third derivatives of $g_{\\mu\\nu}$ and $\\phi$. However, in the case $-2\\kappa_1=\\kappa_2\\equiv\\kappa$ the derivative coupling term reads $\\kappa G_{\\mu\\nu}\\phi^{,\\mu}\\phi^{,\\nu}$ and the order of corresponding field equations is reduced up to second one. Assuming $-2\\kappa_1=\\kappa_2$, we study the spatially-flat Friedman-Robertson-Walker model with a scale factor $a(t)$ and find new exact cosmological solutions. It is shown that properties of the model at early stages crucially depends on the sign of $\\kappa$. For negative $\\kappa$ the model has an initial cosmological singularity, i.e. $a(t)\\sim (t-t_i)^{2/3}$ in the limit $t\\to t_i$; and for positive $\\kappa$ the universe at early stages has the quasi-de Sitter behavior, i.e. $a(t)\\sim e^{Ht}$ in the limit $t\\to-\\infty$, where $H=(3\\sqrt{\\kappa})^{-1}$. The corresponding scalar field $\\phi$ is exponentially growing at $t\\to-\\infty$, i.e. $\\phi(t)\\sim e^{-t/\\sqrt{\\kappa}}$. At late stages the universe evolution does not depend on $\\kappa$ at all; namely, for any $\\kappa$ one has $a(t)\\sim t^{1/3}$ at $t\\to\\infty$. Summarizing, we conclude that a cosmological model with nonminimal derivative coupling of the form $\\kappa G_{\\mu\\nu}\\phi^{,\\mu}\\phi^{,\\nu}$ is able to explain in a unique manner both a quasi-de Sitter phase and an exit from it without any fine-tuned potential. ",
        "introduction": "} For many years scalar fields have been an object of great interest for physicists. The reasons for this are manifold. One of them is quite pragmatic: models with scalar fields are relatively simple, and therefore it appeared possible to study them in detail and then extrapolate the results to more realistic and complicated models. More physical motivations include such important topics as the idea about variable ``fundamental'' constants, the Jordan-Brans-Dicke theory initially elaborated to solve the Mach problem, the Kaluza-Klein compactification scheme, the low-energy limit of the superstring theory, and others. Scalar fields play an especially important role in cosmology. As a bright example, one may mention numerous inflationary models in which inflation in the early Universe is typically driven by a fundamental scalar field called an inflaton. Furthermore, a recent discovery of cosmic acceleration has only refreshed the interest to scalar fields which began to be considered as candidates to explain dark energy phenomena.  The rather general form of action for a scalar-tensor theory of gravity with a single scalar field $\\phi$ can be given as\\footnote{Throughout this paper we use units such that $G=c=1$. The metric signature is $(- + + +)$ and the conventions for curvature tensors are $R^\\alpha_{\\beta\\gamma\\delta} = \\Gamma^\\alpha_{\\beta\\delta,\\gamma} - ...$ and $ R_{\\mu\\nu} = R^\\alpha_{\\mu\\alpha\\nu}$.} \\beq\\label{STTaction} S=\\int d^4x\\sqrt{-g}\\left\\{\\frac1{8\\pi}F(\\phi,R) -h(\\phi)g_{\\mu\\nu}\\phi^{,\\mu}\\phi^{,\\nu}\\right\\}, \\eeq where $g_{\\mu\\nu}$ is a metric, $g=\\det(g_{\\mu\\nu})$, and $R$ is the scalar curvature. Functions $F(\\phi,R)$ and $h(\\phi)$ are varying from theory to theory. The function $h(\\phi)$ is responsible for the sign of kinetic energy of the scalar field. For example, the choice $h(\\phi)\\equiv-1$ leads to a wide class of theories with the negative kinetic term. The function $F(\\phi,R)$, being in general nonlinear, provides a nonminimal coupling between a scalar field and curvature. Though a freedom in choosing of $F(\\phi,R)$ leads to an unlimited variety of scalar-tensor theories, it is known (see, for example, \\cite{Mae,FarGunNar,CroFra}) that there exist conformal transformations transforming this kind of theories to Einstein's theory with a new minimally coupled scalar field $\\phi$ and an effective potential $V(\\phi)$ describing its self-interaction. The potential $V(\\phi)$ is a very important ingredient of various cosmological models: a slowly varying potential behaves like an effective cosmological constat providing one or more than one inflationary phases. An appropriate choice of $V(\\phi)$ is known as a problem of fine tuning of the cosmological constant.  A further extension of scalar-tensor theories can be represented by models with nonminimal couplings between derivatives of a scalar field and curvature. This kind of couplings may appear in some Kaluza-Klein theories \\cite{KK1,KK2} (see also \\cite{Lindebook}, Section 9.5). In 1993, Amendola \\cite{Ame} has been considered the most general gravity Lagrangian linear in the curvature scalar $R$, quadratic in $\\phi$, and containing terms with four derivatives including all of the following terms (see also \\cite{CapLamSch} for details): \\bea & L_1=\\kappa_1 R\\phi_{,\\mu}\\phi^{,\\mu};\\quad% L_2=\\kappa_2 R_{\\mu\\nu}\\phi^{,\\mu}\\phi^{,\\nu};\\quad% L_3=\\kappa_3 R \\phi\\square\\phi ; &\\nonumber\\\\ & L_4=\\kappa_4 R_{\\mu\\nu} \\phi\\phi^{;\\mu\\nu} ;\\quad% L_5=\\kappa_5 R_{;\\mu} \\phi\\phi^{,\\mu} ;\\quad% L_6=\\kappa_6 \\square R \\phi^2, &\\nonumber \\eea where coefficients $\\kappa_1,\\dots,\\kappa_6$ are coupling parameters with dimensions of length-squared. Using the divergencies $$ (R\\phi^{,\\mu}\\phi)_{;\\mu};\\quad (R^{\\mu\\nu}\\phi\\phi_{,\\mu})_{;\\nu};\\quad% (R^{;\\mu}\\phi^2)_{;\\mu}, $$ one may conclude that, without loss of generality, $L_4$, $L_5$, and $L_6$ are not necessary to be considered. Also one may rule out $L_3$ because it contains $\\phi$ itself, while coupling term of main interest are those, where only the gradient of $\\phi$ is included. Thus, a general scalar-tensor theory with nonminimal derivative couplings may include only two terms $L_1$ and $L_2$.  As was shown by Amendola \\cite{Ame}, a theory with derivative couplings cannot be recasting into Einsteinian form by a conformal rescaling $\\tilde g_{\\mu\\nu}=e^{2\\omega}g_{\\mu\\nu}$. He also supposed that an effective cosmological constant, and then the inflationary phase can be recovered without considering any effective potential if a nonminimal derivative coupling is introduced. Amendola himself \\cite{Ame} has considered a cosmological model in the theory with the only derivative coupling term $L_1=\\kappa_1 R\\phi_{,\\mu}\\phi^{,\\mu}$ and, by using a generalized slow-rolling approximation (i.e., neglecting all terms of order higher than the second one), he has obtained some analytical inflationary solutions. A general model containing both $L_1$ and $L_2$ has been discussed in \\cite{CapLamSch} (see also \\cite{CapLam}); it was shown that the de Sitter spacetime is an attractor solution of the model if $4\\kappa_1+\\kappa_2>0$. Recently Daniel and Caldwell \\cite{DanCal} have considered a theory with the derivative coupling term $L_2=\\kappa_2 R_{\\mu\\nu}\\phi^{,\\mu}\\phi^{,\\nu}$; in particular, they studied constraints which precision tests of general relativity impose on the coupling parameter $\\kappa_2$. It is also worth mentioning a series of papers devoted to a nonminimal modification of the Einstein-Yang-Mills-Higgs theory \\cite{BalDehZay:07} (see also a review \\cite{BalDehZay} and references therein).  In this paper we continue studying a scalar-tensor theory with nonminimal derivative couplings and construct new exact cosmological solutions of the theory. ",
        "conclusions": "We have considered the gravitational theory of a scalar field with nonminimal derivative coupling to curvature and studied cosmological models in this theory. The main results obtained are as follows:  1. The Lagrangian of the theory includes two derivative coupling terms $\\kappa_1 R\\phi_{,\\mu}\\phi^{,\\mu}$ and $\\kappa_2 R_{\\mu\\nu}\\phi^{,\\mu}\\phi^{,\\nu}$, where $\\kappa_1$ and $\\kappa_2$ are coupling parameters with dimensions of length-squared. In general, field equations of the theory are of third order, i.e., contain third derivatives of $g_{\\mu\\nu}$ and $\\phi$, but in the particular case the order of equations is reduced up to the second one. This case corresponds to the choice $-2\\kappa_1=\\kappa_2\\equiv\\kappa$, then a combination of derivative coupling terms turn into $\\kappa G_{\\mu\\nu}\\phi^{,\\mu}\\phi^{,\\nu}$. It is worth noting that Capozziello et al \\cite{CapLamSch} , at pages 43 and 47, have mentioned the case $-2\\kappa_1=\\kappa_2$ to play a special role, because it represents a singular point of the differential equation. In this paper, we have supposed that the theory with $-2\\kappa_1=\\kappa_2$ is more preferable with the physical point of view, since the corresponding field equations do not contain derivatives of dynamical variables of order higher than the second.  2. Assuming $-2\\kappa_1=\\kappa_2\\equiv\\kappa$, we have studied a cosmological model with the spatially-flat Friedman-Robertson-Walker metric. It was shown that a behavior of the scale factor $a(t)$ and the scalar field $\\phi$ at large times is the same for all values of $\\kappa$ including zero, that is the late evolution of universe does not depend on $\\kappa$. Namely, one has $a(t)\\sim t^{1/3}$ and $\\phi(t)\\sim\\ln t$ at $t\\to\\infty$. Note this asymptotical behavior coincides with that of the exact solution \\Ref{alpha0}, \\Ref{phi0} obtained for $\\kappa=0$ (no coupling).  3. General properties of the model crucially depends on a sign of $\\kappa$. For $\\kappa<0$ an asymptotical form of the cosmological metric for small times is given by Eq. \\Ref{metric_k<0}. A corresponding scale factor is $a(t)\\sim (t-t_i)^{2/3}$; it describes the universe with an initial singularity at $t=t_i$. A new interesting feature of the model with derivative coupling is that a behavior of the scalar field near the cosmological singularity is regular, $\\phi(t)\\sim t$ (see Eq. \\Ref{phi_k<0}). For $\\kappa>0$ the law of universe evolution is qualitatively distinct from that for $\\kappa<0$. Now at early stages the universe has the quasi-de Sitter behavior \\Ref{deSitter} corresponding to the cosmological constant $\\Lambda=(3\\kappa)^{-1}$. In the limit $t\\to-\\infty$ the scale factor has the following asymptotical form $a(t)\\sim \\exp\\left(\\frac{t}{3\\kappa^{1/2}}-\\frac{e^{2t/\\sqrt{\\kappa}}}{144\\pi\\kappa C^2}\\right)$ (see Eq. \\Ref{a_k>0}), hence $a(t)$ exponentially fast goes to the de-Sitter form $a(t)=e^{Ht}$ with $H=(3\\sqrt{\\kappa})^{-1}$. At the same time, the scalar field $\\phi$ is exponentially growing at $t\\to-\\infty$, namely $\\phi(t)\\sim e^{-t/\\sqrt{\\kappa}}$.   In conclusion, let us summarize the most essential features of cosmology with nonminimal derivative coupling. First of all, we should emphasize that cosmological solutions with the quasi-de Sitter phase are typical solutions of the gravitational theory of a scalar field with derivative coupling of the form $\\kappa G_{\\mu\\nu}\\phi^{,\\mu}\\phi^{,\\nu}$ with positive $\\kappa$. So, in order to obtain an inflationary phase, one need no fine-tuned potential, and so one do not face with the problem of fine-tuning. Another important feature of the model consists in the fact that an exact cosmological solution with $\\kappa>0$ describes in a unique manner both a quasi-de Sitter phase and an exit from it. Thus, the problem of graceful exit from inflation in cosmology with the derivative coupling term $\\kappa G_{\\mu\\nu}\\phi^{,\\mu}\\phi^{,\\nu}$ has a natural solution without any fine-tuned potential.      \\subsection*"
    },
    "0910/0910.2396_arXiv.txt": {
        "abstract": "Recent surveys have revealed a lack of close-in planets around evolved stars more massive than 1.2~$M_{\\odot}$. Such planets are common around solar-mass stars. We have calculated the orbital evolution of planets around stars with a range of initial masses, and have shown how planetary orbits are affected by the evolution of the stars all the way to the tip of the Red Giant Branch (RGB). We find that tidal interaction can lead to the engulfment of close-in planets by evolved stars. The engulfment is more efficient for more-massive planets and less-massive stars. These results may explain the observed semi-major axis distribution of planets around evolved stars with masses larger than 1.5~$M_{\\odot}$.  Our results also suggest that massive planets may form more efficiently around intermediate-mass stars. ",
        "introduction": "Observationally, due mostly to the inapplicability of high-precision Doppler techniques to stars with spectral types earlier than late-F (which have large rotational velocities and a small number of spectral lines), little was known until recently about the frequency of planets around the more-massive stars.  This situation has now changed, as surveys have been extended to  searches for planets around more-massive stars in an evolved stage \\citep{Doli07,Doli09,Frin02,Hat03,Hat05,Hat06,Joh07a,Joh07b,Joh08,Liu07,Lm07, Nied07,Nied09,Ref06,Sat07,Sat08,Set03,Set05}.  One of the most important trends that these surveys have revealed is the lack of close-in planets orbiting stars with masses $M > 1.5~M_{\\sun}$ \\citep{Joh07b,Sat08,Wri09} despite the fact that these planets are found around $\\approx$20 \\% of the Main Sequence (MS) stars with $M < 1.2~M_{\\sun}$. The frequency of planets seems also to be higher around intermediate-mass stars \\citep{Lm07,Joh07b}. Furthermore, contrary to the correlation between metallicity and probability of planet-hosting found for solar-mass, main-sequence stars \\citep[e.g.,][]{Fv05,San05}, the more-massive planet-hosting stars do not exhibit higher metallicities \\citep[e.g.,][]{Pas07}.  On the theoretical side, it has been shown that the formation of Jupiter-mass planets around M~stars may be hindered \\citep[e.g.,][]{L04,Il05}, while the probability that a given star has at least one gas giant increases linearly with the stellar mass up of 3~$M_{\\odot}$ \\citep[e.g.,][]{Kk08}.  One possibility is that the observed difference in the orbital distribution of planets found around intermediate and solar-mass primaries is due to the evolution of the star. Since all of the planets orbiting red giants or subgiants with $M > 1.5~M_{\\sun}$ have semi-major axes $a > 0.5$~AU it has been suggested that the planets might be engulfed as the star evolves off the MS \\citep{Joh07b,Sat08}. This possibility is, however, often dismissed out in the literature with the argument that high-mass stars are physically too small to engulf hot Jupiters \\citep{Joh07b,Joh08,Cur09}. Another possibility is that the observed differences in orbital distribution are primordial, and and they are a consequence of the planet formation mechanism around more massive stars.  Along these lines, it has been shown by \\cite{Cur09} that the dependence of the lifetime of the gaseous disk on the stellar mass could result in halting the inward migration of planets around high-mass stars,  thus explaining the observed lack of short period planets around these stars.  The point we are making in the present paper is that before stellar evolution can be ruled out as the mechanism behind the observed semi-major axis distribution of planets around evolved stars, detailed modeling of the orbital evolution needs to be performed. In other words, in order to determine the potential role of the stellar mass in the planet formation process, the effects of the evolution of the star on the observed orbit distribution around giants have to be correctly isolated. This is precisely the goal of this paper. ",
        "conclusions": "Our main goal has been to determine whether stellar evolution could explain the observed distribution of the semi-major axes of planetary orbits  around evolved stars (i.e., semi-major axis $>$\\,0.5~AU). We found that when the details of the orbital evolution are accurately calculated, tidal interactions constitute a quite powerful mechanism, capable of capturing close-in planets into the envelope of evolved stars.  To date, there are $\\sim$20 exoplanets discovered around giant stars with $M > 1.5~M_{\\sun}$. The host stars have radii in the range $0.02 < R_* < 0.1$~AU. Our models are consistent with the existence of these planets at the point in the RGB evolution at which they are observed (see e.g., Table~8 in \\citealt{Sat08} and our Figs.~\\ref{fig_2m} and~\\ref{fig_3m}).  However, we do not expect to find massive planets with a $<$\\,0.4~AU around a 2~$M_\\sun$ star with a R$_*\\geq 0.1$~AU (or 24~$R_\\sun$). Our calculations  provide the minimum orbital radius inside of which planets will be engulfed by the star at the end of the RGB evolution.  We find that the evolution of the star alone can quantitatively explain the observed lack of close-in planets around evolved stars even allowing for the uncertainties associated with mechanisms such as mass loss along the RGB or tidal-interaction theory. A mechanism such as the one invoked by Currie (2009; i.e., a lifetime stellar-mass dependency of the gas in the planet-forming disk that can halt migration) is not needed, although it might still be present.  We find that given an initial distance at which tidal capture is possible, the more massive the planet, the earlier it will be captured by the RGB envelope. Observationally, it appears that giant stars host more massive planets than MS stars \\citep[e.g.,][]{Joh07b,Lm07}. Since we find that more massive planets are expected to be engulfed earlier in the RGB evolution, the fact that they are more frequently observed may point towards a planet-formation mechanism that favors the formation of more-massive planets around intermediate-mass stars. This would be consistent with a scenario in which the disk mass scales with the stellar mass, and more massive disks produce more massive planets (see also \\citealt{Kk08}).  Along similar lines, since we find a high probability of tidal capture of the planet by evolved stars, the higher frequency of planets observed around intermediate-mass stars \\citep{Lm07,Joh07b} seems to imply that the efficiency of planet formation must be considerably higher for more massive stars, compared to their solar analogous.  \\cite{Ass09} estimated the probability of detecting transits of planetary companions to giant stars to be $\\geq$ 10 \\% for several of the known systems. Since tidal oribital decay decreases the initial orbital distance, it may increase the probability for the planet to be observed in transit.  Our results suggest that many planets may be accreted by their host star \\citep[see also][]{SL99a,SL99b}. Although so far the results are based on a fairly limited sample, it appears that giant stars hosting planets have the same metallicity distribution as giant stars without planets \\citep{Pas07}. On the other hand, we should note that if a giant star engulfs a close-in planet early in the RGB, it will appear to be a giant star without a planet. This mechanism might perhaps provide for a partial explanation for the lack of correlation with metallicity.  Finally, stars that have accreted planets may show a higher spin rate, as the planet's orbital angular momentum is transferred to the star \\citep[e.g.][]{LS02}. This phenomenon has been investigated by \\citet{Mass08} and \\citet{Carl09}, where they estimate the probability of finding rapid rotators among evolved stars. Our quantification of the tidal capture radius may help refine these calculations."
    },
    "0910/0910.2675_arXiv.txt": {
        "abstract": "We report a moderate-depth (70 ksec), contiguous 0.7  \\sqdeg , \\chandra\\ survey, in the Lockman Hole Field of the {\\it Spitzer}/SWIRE Legacy Survey coincident with a completed, ultra-deep VLA survey with deep optical and near-infrared (NIR) imaging in-hand.  The primary motivation is to distinguish starburst galaxies and Active Galactic Nuclei, including the significant, highly obscured (log \\nh $>$23) subset. \\chandra\\ has detected 775 X-ray sources to a limiting broad band (0.3-8 keV) flux $\\sim 4 \\times 10^{-16}$ \\fcgs . We present the X-ray catalog, fluxes, hardness ratios and multi-wavelength fluxes. The log N vs. log S agrees with those of previous surveys covering similar flux ranges. The Chandra and {\\it Spitzer} flux limits are well matched: 771 (99\\%) of the X-ray sources have IR (infared) or optical counterparts, and 333 have MIPS 24 $\\mu$m detections. There are 4 optical-only X-ray sources and 4 with no visible optical/IR counterpart. The very deep ($\\sim$2.7 $\\mu$Jy rms) VLA data yields 251 ($> 4 \\sigma$) radio counterparts, 44\\% of the X-ray sources in the field. We confirm that the tendency for lower X-ray flux sources to be harder is primarily due to absorption. As expected, there is no correlation between observed IR and X-ray flux. Optically bright, Type 1 and red AGN lie in distinct regions of the IR vs X-ray flux plots, demonstrating the wide range of SEDs in this sample and providing the potential for classification/source selection. Many optically-bright sources, which lie outside the AGN region in the optical vs X-ray plots ($\\rm f_r / f_x > 10$), lie inside the region predicted for red AGN in IR vs X-ray plots consistent with the presence of an active nucleus. More than 40\\% of the X-ray sources in the VLA field are radio-loud using the classical definition, \\rl . The majority of these are red and relatively faint in the optical so that the use of \\rl~to select those AGN with the strongest radio emission becomes questionable. Using the 24 $\\mu$m to radio flux ratio (\\q24 ) instead results in 13 of the 147 AGN with sufficient data being classified as radio-loud, in good agreement with the $\\sim$10\\% expected for broad-lined AGN based on optical surveys. We conclude that \\q24 ~is a more reliable indicator of radio-loudness. Use of \\rl ~should be confined to optically-selected, Type 1 AGN. ",
        "introduction": "SWIRE, the largest {\\it Spitzer} Legacy program, is tracing the evolution of dusty, star-forming galaxies, evolved stellar populations and AGN, as a function of environment from z \\gax 2.5 to the present epoch.  SWIRE covers 6 fields with a total area of 49 \\sqdeg\\ in all seven {\\it Spitzer} bands, selected from the entire IRAS/DIRBE sky as those areas with the lowest 100$\\mu m$ surface brightness (which scales with $N_H$, Schlegel \\etal\\ 1998); the Lockman Hole is one of the two best regions, having several contiguous \\sqdeg\\ of low $N_H$($\\sim 7 \\times 10^{19}cm^{-2}$) sky.  SWIRE's power comes from its large surface area, its depth, which probes the Universe out to redshifts z \\gax 2.5 and over time intervals $> 10 Gyr$, and its sensitivity to both evolved stellar populations (IRAC: 3--8$\\mu m$) and dusty objects (starbursts \\& obscured AGN; MIPS: 24,70,160$\\mu m$). Thus their distribution relative to environment, and evolution with respect to time and to the development of structure, can be studied together. SWIRE is one of several, on-going, medium-depth surveys which fill the gap between the deep and shallow surveys, dramatically illustrating the need for a multi-layered approach to efficiently fill the L-z plane in several wave-bands.  The SWIRE prime science goal is to study the structure, evolution and environments of AGN, starbursts and spheroidal galaxies out to z\\gax 2.5. Key to this extensive multi-wavelength campaign is an X-ray survey deep enough to distinguish starbursts and AGN, including the significant numbers which are highly obscured \\nh\\gax 10$^{23}$ cm$^{-2}$. We carried out a \\chandra\\  X-ray survey  in the best (in terms of Galactic extinction and absence of nearby bright sources, including no bright radio sources) extragalactic $\\sim$1 \\sqdeg\\ field within SWIRE. The X-ray observations cover 0.7 \\sqdeg\\ contiguously. The broad band (0.3-8.0 keV) flux limit is $4 \\times 10^{-16}$ \\fcgs\\ at the field centers, sufficient to detect all SWIRE IR AGN except for those with high absorption at low redshift (z$\\leq$0.8, log N$_H \\geq$ 24) while IR galaxies are not detected. Thus we are able to distinguish all but the most highly obscured AGN from amongst the IR sources by their X-ray emission.  The standard AGN model includes a super-massive black hole surrounded by an accretion disk in the center of a galaxy. The central regions of the AGN produce strong, hard X-ray emission (and relativistic jets in radio-loud sources) along with thermal emission from an accretion disk. Gas in the vicinity is heated by the nuclear emission and produces the broad and narrow emission lines characteristic of Type 1 and Type 2 AGN respectively. The viewing orientation of an AGN effects its classification, as demonstrated by the detection of polarized broad lines in NGC1068 (Antonucci \\& Miller 1985). This result led to a general acceptance that some fraction of Type 2 AGN are edge-on Type 1s. Such unification schemes require highly opaque dust surrounding the accretion disk. In addition to obscuring our view to the nucleus, this dust is heated by the nuclear source producing strong IR emission. Thus the observed Spectral Energy Distribution (SED) of an AGN is dependent on viewing angle and the traditional optical/ultraviolet surveys are incomplete to obscured sources ({\\it e.g.} Polletta \\etal\\ 2006). It is now clear that there is a significant obscured AGN population which has been missed from earlier surveys. The powerful combination of radio, X-ray and IR observations available for the current surveys facilitates a different and potentially unbiased view of the AGN in the field and so probes the full AGN population, including obscured sources. The variety of new techniques has resulted in new types of AGN being found, including: red AGN (Cutri \\etal\\ 2002), X-ray bright optically-normal galaxies (XBONGS, Comastri et al. 2002), Type 2 quasars (Norman et al. 2002), X-ray detected Extremely Red Objects (V/EROs, Alexander \\etal\\ 2002, Brusa \\etal\\ 2005).  Despite the several multi-wavelength surveys in progress during this era of Great Observatories, a full view of the AGN population remains elusive. Shallow surveys, which are dominated by Type 1 AGN, find largely distinct subsets of the population depending on waveband (Hickox et al. 2009). While overlap is significant in deeper surveys, selection in any single waveband does not find all the AGN (Barmby et al 2006, Polletta et al. 2006, Donley et al. 2008, Park et al. 2008). Samples of obscured AGN remain relatively small and biases are still in the process of being understood, so that the nature of the new AGN and their significance to the population as a whole remain undetermined. Compton thick AGN are hard to find because their X-ray flux is obscured to energies \\gax 10 keV. Estimates of the fraction of the general population which are Compton thick vary. X-ray selected AGN include $\\sim 30$\\% (Polletta \\etal\\ 2006, Treister et al. 2004) Compton-thick sources. This number doubles to $\\sim 66\\%$ when IR-selected AGN are included (Polletta et al. 2006, Treister et al. 2009). Estimates generated by stacking X-rays from IR-selected AGN also suggest a factor of two increase in IR-selected compared with X-ray selected Compton thick AGN (Fiore et al. 2008, Daddi et al. 2007), although the absolute numbers differ by factors of $\\sim$100 depending on selection techniques and assumptions. The latest estimate from modeling the Cosmic X-ray Background (CXRB) postulates an obscured population larger than the unobscured by a factor of 8 at low luminosities, half of which are Compton thick. This decreases to a factor of 2 at high luminosities (Gilli, Comastri \\& Hasinger 2007). Alternatively, many X-ray-selected, low-luminosity AGN may be intrinsically X-ray hard rather than obscured, possibly reducing the fraction of obscured (including Compton thick) AGN to $\\sim$20\\% (Hopkins et al. 2009).    A number of wide area, multi-wavelength surveys with sufficient depth to view the full AGN population are currently in progress, including: ECDFS (Lehmer et al. 2005, 0.3 deg$^2$), OPTX (Trouille et al. 2008, 1 deg$^2$ non-contiguous), SWIRE/\\chandra\\ (this paper, 0.7 deg$^2$), AEGIS-X (Laird et al. 2009, 0.67 deg$^2$), and C-COSMOS (Elvis et al. 2009, 0.5+0.4 deg$^2$). They are beginning to provide significant samples of the relatively rare sources, such as the high redshift or obscured AGN, found in small numbers in the deep \\chandra\\ surveys (Luo et al. 2008, Alexander et al. 2003). These various surveys will allow us to properly characterize these populations and understand their significance to AGN as a whole.  In this paper we present the 0.7 \\sqdeg~SWIRE/\\chandra\\ X-ray source catalog, SWIRE identifications and IR, optical and radio fluxes. IR properties of X-ray sources range from the power-law shape characteristic of AGN-dominated sources, mostly unobscured Type 1 AGN, to SEDs which are dominated by star formation or host-galaxy emission, implying obscuration of the AGN at these wavelengths (Franceschini et al. 2005, Barmby et al. 2006, Polletta et al. 2007, Feruglio et al. 2008, Cardamone et al. 2008, Gorjian et al. 2008). X-ray hardness is loosely related to the optical/IR colors, generally supporting this view. We characterize the X-ray and multi-wavelength properties of the SWIRE/\\chandra\\ sample and take an initial look at these properties as a function of X-ray hardness and radio-loudness. Detailed study and modeling of the spectral energy distributions (SEDs) will be presented in a companion paper (Polletta \\etal\\, in prep.). ",
        "conclusions": "We present a list of 775 \\chandra\\ X-ray sources in the SWIRE/\\chandra\\ medium-depth, X-ray survey. Cross-correlation with {\\it Spitzer}, optical and radio images of part/all of the field resulted in 771 (99\\%) identifications in at least one optical/IR band, 767 have IR counterparts visible in the {\\it Spitzer} data in at least one (IRAC) band and 333 have 24$\\mu$m with MIPS detections. Four of the sources have no optical/IR counterpart down to our flux limits and 4 have an optical but no IR couterpart. We present multi-wavelength flux measurements for 744 X-ray sources, all those which are uncontaminated, unconfused and above the formal survey thresholds. The near-IR$-$X-ray datasets are well-matched in flux limit and go deep into the AGN population, providing an excellent dataset for multi-wavelength studies of the full AGN population, which will be reported in a companion paper (Polletta et al., in prep.). As in earlier surveys (DK04), there is no correlation between X-ray hardness and hard X-ray flux in this sample, confirming that hardness is predominantly caused by obscuration in the X-rays.  The very deep (2.7 $\\mu$Jy at the field center) VLA data, covering part of the \\chandra\\ field (Figure~\\ref{fg:image}), results in 251 ($> 4 \\sigma$) radio detections, $44$\\% of the 568 X-ray sources in the VLA field. Even the very deepest radio data cannot detect all the X-ray sources. We demonstrate that the traditional radio-to-optical flux ratio, \\rl, used to define radio-loudness in AGN breaks down for a large proportion of the sources in this X-ray selected sample due to the weakness of the optical emission. Use of the 24$\\mu$m flux in place of the optical, the \\q24 ~parameter (Appleton et al. 2004), brings the radio-loud fraction down to expected levels (9\\%) and is strongly preferred as an indicator of radio-loudness for the full AGN population.   The wide range of optical/IR/X-ray SEDs in this X-ray selected sample is demonstrated by the lack of any correlation between the optical or IR flux and the X-ray flux. Comparison with tracks for optically-selected, Type 1, and red AGN SEDs demonstrates the predominance of red AGN in this sample. There is a continuous distribution rather than distinct classes of AGN, and the IR vs X-ray plots allow better discrimination than the optical as the effects of obscuration are much lower. Comparison of the source properties (X-ray hardness, flux ratios and radio loudness) with the predictions of the standard median SEDs allows us to broadly classify the source types as a function of position in flux-flux plots. A correlation between the 24$\\mu$m flux and the X-ray hardness disappears when the hard X-ray flux is used, demonstrating that it is due to absorption in the X-rays while the 24$\\mu$m flux remains uneffected. This reinforces the use of \\q24 ~to define radio-loudness due to its stability."
    },
    "0910/0910.5325_arXiv.txt": {
        "abstract": "% During the period July 2007 - January 2009, the AGILE satellite, together with several other space- and ground-based observatories monitored the activity of the flat-spectrum radio quasar 3C~454.3, yielding the longest multiwavelength coverage of this \\gray quasar so far. The source underwent an unprecedented period of very high activity above 100 MeV, a few times reaching \\gray flux levels on a day time scale higher than $F=400 \\times 10^{-8}$\\,\\phcmsec, in conjunction with an extremely variable behavior in the optical $R$-band, even of the order of several tenth of magnitude in few hours, as shown by the GASP-WEBT light curves. We present the results of this long term multiwavelength monitoring campaign, with particular emphasis on the study of possible lags between the different wavebands, and the results of the modeling of simultaneous spectral energy distributions at different levels of activity. ",
        "introduction": "% Among active galactic nuclei (AGNs), blazars show intense and variable \\gray emission above 100~MeV \\citep{Hartman1999:3eg}, and the variability time scale can be as short as a few days, or last a few weeks. They emit across several decades of energy, from the radio to the TeV energy band and their spectral energy distributions (SEDs) are typically double humped with a first peak occurring in the IR/optical band in the flat-spectrum radio quasars (FSRQs) and low-energy peaked BL Lacs (LBLs), and at UV/X-rays in the high-energy peaked BL Lacs (HBLs). This peak is commonly interpreted as synchrotron radiation from high-energy electrons in a relativistic jet. The second SED component is commonly interpreted as inverse Compton (IC) scattering of soft seed photons by relativistic electrons, and peaks in the MeV--GeV and in the TeV energy bands in the FSRQs/LBLs and in the HBLs, respectively. A recent review of the blazar emission mechanisms and energetics is given in \\cite{Celotti2008:blazar:jet}.  The FSRQ \\source{} (PKS~2251$+$158; $z=0.859$) is certainly one of the most active extragalatic sources at high energy. In the \\egret{} era, it was detected in 1992 during an intense \\gray flaring episode \\citep{Hartman1992:3C454iauc, Hartman1993:3C454_EGRET} when its flux $F_{\\rm E>100MeV}$ was observed to vary within the range $(0.4-1.4) \\times 10^{-6}$\\,photons\\,cm$^{-2}$\\,s$^{-1}$. In 1995, a 2-week campaign detected a \\gray flux $< 1/5$ of its historical maximum \\citep{Aller1997:3C454_EGRET}.  Figure~\\ref{Fig:3c454:egret} shows the \\gray light curve for $E>100$\\,MeV as observed by EGRET in the period 1991--1995. \\begin{figure}[!ht] \\begin{center} \\includegraphics[width=8cm]{vercellones_f1.eps} \\end{center} \\vspace{-0.5truecm} \\caption{EGRET \\source{} light curve for $E>100$\\,MeV in the period 1991-1995. The downward arrow represents a $2\\sigma$ upper-limit. Data from \\cite{Hartman1999:3eg}.} \\label{Fig:3c454:egret} \\end{figure}   In 2005, \\source{} underwent a  major flaring activity in almost all energy bands (see \\citealt{Giommi2006:3C454_Swift}). In the optical, it reached $R=12.0$\\,mag \\citep{vil06} and it was detected by \\igr{} at a flux\\footnote{Assuming a Crab-like spectrum.} level of $\\sim 3 \\times 10^{-2}$\\,photons\\,cm$^{-2}$\\,s$^{-1}$ in the 3--200~keV energy band \\citep{Pian2006:3C454_Integral}. Since the detection of the exceptional 2005 outburst, several monitoring campaigns were carried out to follow the source multifrequency behavior \\citep{vil06,vil07,rai07,rai08a,rai08b}. During the last of these campaigns, 3C~454.3 underwent a new optical brightening in mid July 2007, which triggered observations at all frequencies, including the \\agile{} one. In the following, we briefly describe the \\agile{} satellite, and we discuss the various campaigns triggered on \\source{}. ",
        "conclusions": "Since July 2007, \\source{} has been playing the same role for \\agile{} as 3C~279 had for EGRET, and during  the period July 2007 - January 2009 we acquired data not only in the \\gray energy band, but across 14 decades in energy. This allowed us to construct simultaneous SEDs, sampling high, intermediate, and low \\gray emission states, involving both ground and space based observatories.  We found that the role of the external Compton on the disk and the broad-line region radiation (and possibly also on hot corona photons) is crucial to account for the hard \\gray spectrum states. Moreover, thanks to the extremely dense optical coverage provided by the GASP-WEBT, we were able to study the correlations between the \\gray and optical fluxes. We found a $\\simeq 1$ day possible lag of the high energy photons with respect to the optical ones.  The simultaneous presence of two \\gray satellites, \\agile{} and Fermi, the extremely prompt response of wide-band satellites as \\swi{}, and the long-term monitoring provided from the radio to the optical by the GASP-WEBT Consortium will assure the chance to investigate and study the physical properties of \\source{} and of several more blazars both at high and low emission states."
    },
    "0910/0910.4323_arXiv.txt": {
        "abstract": "As of today over 40 planetary systems have been discovered in binary  \\langeditorchanges{star systems}. In all cases the configuration appears to be circumstellar, where the planets orbit around one of the stars, the secondary acting as a perturber.  The formation of planets in binary \\langeditorchanges{star systems} is more difficult than around single stars due to the gravitational action of the companion on the dynamics of the protoplanetary disk. In this contribution we first  \\langeditorchanges{briefly} present the relevant observational evidence  \\langeditorchanges{for} planets in binary systems. Then the dynamical influence that a secondary companion has on a circumstellar disk will  \\langeditorchanges{be} analyzed through \\langeditorchanges{fully} hydrodynamical simulations. We demonstrate that the disk becomes eccentric and shows a coherent precession around the primary star. Finally, fully hydrodynamical simulations of evolving protoplanets embedded in disks  in binary \\langeditorchanges{star systems} are presented. We investigate how the orbital evolution of protoplanetary embryos and their mass growth from cores to massive planets might be affected in this very dynamical environment. We  \\langeditorchanges{consider, in particular,} the planet orbiting the primary in the system $\\gamma$ Cephei. ",
        "introduction": "Planet formation is obviously a process that occurs around single as well as in multiple \\langeditorchanges{star systems}, a fact that is indicated by the detection of well over 40 planetary systems that reside in a binary or even multiple star configurations. All of the observed systems display a so called S-type configuration in which the planets orbit around one of the stars and the additional star, the companion or secondary star, acts as a perturber to this system. In this review we shall refer to the secondaries as single objects, even though they may be multiple. As indicated in Table~\\ref{kley-tab:systems} the distances of the secondaries \\langeditorchanges{from} the host \\langeditorchanges{stars} of the planetary \\langeditorchanges{systems} range  from very small values of about 20 AU for Gl~86 and $\\gamma$ Cep to several thousand AU. There are now 4 confirmed systems with a binary separation in the order of 20 AU; in \\langeditorchanges{Table~}\\ref{kley-tab:systems} \\langeditorchanges{these are} shown below the horizontal separation line. The mere existence of these 4 \\langeditorchanges{systems represents} a special challenge to any kind of planet formation process, due their tightness. Interestingly, there appears to be a lack of planets for \\langeditorchanges{intermediate separations} as there are no planets in binaries with separations between 20 and 100 AU. There are many more systems with larger \\langeditorchanges{separations} (not listed in the table), but in most of the cases only projected distances can be given, and the real physical \\langeditorchanges{separations are necessarily} larger.  \\begin{table}[h,t] \\label{kley-tab:systems} \\begin{center} \\begin{tabular}{|l|l|l|l|l|l|} \\hline Star  & $a_\\mathrm{bin}$ {\\small{[AU]}} &  $a_\\mathrm{p}$  {\\small{[AU]}} & $M\\mathrmpl\\sin$ i  {\\small{ [M$_{\\rm Jup}$]}} &  $e_\\mathrm{p}$ & Remarks \\\\ \\hline HD 40979  & 6400   &  0.811 & ~3.32  & .23 & \\\\ Gl 777 A  & 3000   &  3.65  & ~1.15  & .48 & \\\\ HD 80606  & 1200   &  0.439 & ~3.41  & .93 & \\\\ 55 Cnc B  & 1065   &  0.1-5.9 & ~0.8-4.05  & .02-.34 & {\\small multiple} \\\\ 16 Cyg B  & 850    &  1.66  &  ~1.64  & .63  & \\\\ $\\upsilon$ And  & 750   &  0.06-2.5 &  0.7-4.0  & .01-.27 &  {\\small multiple} \\\\ HD 178911 B  & 640    &  0.32  & ~6.3   & .12  & \\\\ HD 219542 B  & 288    &  0.46  & ~0.30  & .32  & \\\\ $\\tau$ Boo  & 240    &  0.05  &  ~4.08  & .02  & \\\\ HD 195019  & 150    &  0.14  &  ~3.51  & .03  & \\\\ HD 114762  & 130    &  0.35  &  11.03   & .34 &  \\\\ HD 19994  & 100    &  1.54  &  ~1.78  & .33  & \\\\ \\hline HD 41004A  & 23     &  1.33  &  ~2.5   & .39 & {\\small multiple} \\\\ $\\gamma$ Cep  & 20.2   &  2.04  &  1.60   & .11 & {\\small $e_\\mathrm{bin}= 0.4$} \\\\ {HD 196885}  & 17     &  2.63  &  2.96   & .46 & {\\small $e_\\mathrm{bin}= 0.4$} \\\\ Gl 86  & 20     &  0.11  &  4.0    & ~0.046 & {\\small White Dwarf} \\\\ \\hline \\end{tabular} \\end{center} \\caption{Some observed planets in binary \\langeditorchanges{star systems}. This is a selection with emphasis on the shorter period binaries (see \\, \\cite{2004A&A...417..353E,2006ApJ...646..523R,2008A&A...479..271C}). \\langeditorchanges{The} list is very incomplete for larger separations. } \\end{table} Despite the actual detection of planets in binary systems there is additional circumstantial evidence of debris disks (which are thought to be a byproduct of the planet formation process) in binary systems as indicated by Spitzer data. Here, for S-type configurations it is found that disks around an individual star of the binary exist mainly for binary separations larger than 50 AU, while P-type circumbinary debris disks are detected only in very tight binaries with $a_\\mathrm{bin}$ smaller than about 3 AU (\\cite{2007ApJ...658.1289T,2008ApJ...674.1086T,2008arXiv0808.1765Z}).  As first pointed out by \\citet{2004A&A...417..353E}, {see also these proceedings}, there is statistical evidence for two interesting features in the mass-period and eccentricity period distribution of planets residing in binary systems: \\langeditorchanges{planets} with periods smaller than about 40 days tend to have larger masses than their counterparts in single star systems, while at the same time their eccentricities are smaller. This trend has been supported by \\langeditorchanges{the} more recent findings of \\citet{2007A&A...462..345D} who tried to correlate this with the tightness of the \\langeditorchanges{binary,} but the statistics \\langeditorchanges{are still based on small sample sizes} and more data are required.  As the influence of the secondaries on the planet formation process will obviously be smaller for larger distances, we shall focus in this contribution on the more challenging tighter binaries and have used the physical \\langeditorchanges{parameters} of the $\\gamma$ Cep system for our models. Interestingly, $\\gamma$ Cep was one of the very first \\langeditorchanges{stars} which has been suggested to contain an extrasolar planet (of 1.7 $M_\\mathrm{Jup}$): {\\it ``This star has the firmest evidence of a very low mass companion''} \\citep{1988ApJ...331..902C}. A statement unfortunately retracted later by the same team \\citep{1992ApJ...396L..91W}, only to be rediscovered by \\citet{2003ApJ...599.1383H}. Today, this system is one of the tightest binary system known to contain a Jupiter-sized protoplanet. For this reason, it has attracted much attention in past years. Several studies looked at the stability and/or the possibility of (additional) habitable planets in the system \\citep{2004RMxAC..21..222D,2004MSAIS...5..127T,2006ApJ...644..543H,2006MNRAS.368.1599V}. In our studies we have taken the data for $\\gamma$~Cep from \\citet{2003ApJ...599.1383H}. The more recent data by \\citet{2007A&A...462..777N} \\langeditorchanges{only slightly} change the dynamical status. ",
        "conclusions": "In this contribution we have concentrated on the planetary growth process in relatively tight binary stars with particular \\langeditorchanges{attention given} to the system $\\gamma$~Cep. To study the effect of the binary we have followed the evolution of planetary embryos interacting with the ambient protoplanetary \\langeditorchanges{disk,} which is perturbed by the secondary star.  As suspected, the perturbations of the disk, in particular its non-zero eccentricity and the periodic creation of strong tidally induced spiral density arms, lead to non-negligible \\langeditorchanges{effects} on the planetary orbital elements. While embryos placed in the disk at different initial distances from the primary star continue to migrate inwards at approximately the same rate, the eccentricity evolution is markedly different for the \\langeditorchanges{different} cases. If the initial distance is beyond about $a \\gsim 2.7$ AU the eccentricity of the embryo continues to rise to very high \\langeditorchanges{values,} and apparently \\langeditorchanges{the orbit remains bound only due to the damping action of the disk.} The main excitation mechanism of the initial rise of the eccentricity is the perturbed disk and the spiral arms near the outer edge of the disk.  For a disk mass of 3$M_\\mathrm{Jup}$ a $1.6 M_\\mathrm{Jup}$ planet can easily be grown, and the final semi-major axis and eccentricity are also in the observed range of the $\\gamma$~Cep planet for suitable accretion rates onto the planet. One of the major problems in forming a planet in such a close binary system via the core instability model is the problem of the formation of the planetary core in the first place. Due to the large relative velocities induced in a planetesimal \\langeditorchanges{disk,} especially for objects of different \\langeditorchanges{sizes,} the growth process is also problematic in itself.  Hence, the formation of the \\langeditorchanges{Jupiter-sized} planet observed in $\\gamma$~Cep via the standard scenarios remains difficult but may not be impossible. Future research will have to concentrate on additional physical effects such as radiative transport, three-dimensional effects and self-gravity of the disk."
    },
    "0910/0910.3735_arXiv.txt": {
        "abstract": "This paper presents a detailed kinematic and chemical analysis of 16 members of the Kapteyn moving group.  The group does not appear to be chemically homogenous.  However, the kinematics and the chemical abundance patterns seen in 14 of the stars in this group are similar to those observed in the well-studied cluster, $\\omega$ Centauri. Some members of this moving group may be remnants of the tidal debris of $\\omega$ Cen, left in the Galactic disk during the merger event which deposited $\\omega$ Cen into the Milky Way. ",
        "introduction": "The idea of hierarchical galaxy formation, in which galaxies are believed to form from the aggregation of smaller elements (see review by \\citet{Freeman02}), has been around since \\citet{Searle78} first proposed this theory as a challenge to the belief that galaxies formed through the smooth collapse of a large protocloud \\citep{Eggen62}. The identification of debris from these smaller fragments remains of utmost importance in modern studies of theoretical and observational stellar dynamics.  The most massive Galactic globular cluster, $\\omega$ Cen, has several unique physical properties which suggest that there are very significant differences in star formation histories, enrichment processes and structure formation between $\\omega$ Cen and other normal globular clusters \\citep{Bekki03}.  A commonly accepted scenario of formation of $\\omega$ Cen is that it is the surviving nucleus of an ancient dwarf galaxy, the outer envelope of which was entirely removed by tidal stripping as it was accreted by the Galaxy (\\citet{Bekki03} and references therein). Through numerical simulations, \\citet{Bekki03} demonstrated the dynamical feasibility of $\\omega$ Cen forming from an ancient nucleated dwarf galaxy which was accreted into the young Galactic disk.  \\citet{meza05} used numerical simulations to investigate the characteristics of tidal debris from satellite galaxies. They showed that these satellites deposit a large fraction of their stars into either the disc component of the Milky Way or into the halo, showing distinct ``trails\" in the angular momentum -- energy plane, depending on the plane of the satellite's orbit during disruption.  Meza et al.  discussed the presence of the $\\omega$ Cen stellar moving group (i.e. the stellar debris) in two studies of metal-poor stars in the solar neighbourhood, those of \\citet{Beers00} and \\citet{Gratton03}. Figure \\ref{mezafig} (taken from \\citet{meza05}) shows how the $\\omega$ Cen group is distinguished in the angular momentum distribution for the Gratton et al. (2003) sample.  In both this sample and the Beers et al. (2000) sample, the $\\omega$ Cen group appears as an over-density of stars at very specific rotational velocity or angular momentum values.  The radial (U) velocities for the $\\omega$ Cen candidate stars are observed to have a symmetric distribution, ranging from -300 to 300 km s$^{-1}$.  An earlier simulation by \\citet{Bekki03} of the accretion of the $\\omega$ Cen parent galaxy showed a strong plume of debris stars in the solar neighborhood with L$_z$ near $-500$ kpc km s$^{-1}$.  See also the discussion by \\citet{Mizutani03} on the kinematics of tidal debris from the parent galaxy.  \\citet{dinescu02} also provided a theoretical prediction of where the $\\omega$ Cen group would appear kinematically using the metal-poor star sample of Beers et al. (2000).  Figure \\ref{dinescu} shows the angular momentum -- energy plane, in which the shaded zone represents the area where expected $\\omega$ Cen candidate stars would lie.  Dinescu defines this region by assessing three globular clusters ($\\omega$ Cen, NGC 362 and NGC 6779) that are believed to have come from the same original parent galaxy, and argues that their distribution may define the area in which further candidate remnants could lie.  This region covers a large interval of energy over an angular momentum range from about L$_{z}$= -200 to -600 kpc km s$^{-1}$.  We have recently investigated the chemical properties of stars in the Kapteyn moving group, first introduced by Eggen (1962). The stars of the Kapteyn group, as most recently tabulated by Eggen (1996), are mostly metal-poor and in retrograde galactic orbits, so they were identified as a halo moving group.  Moving stellar groups can originate in several ways.  Some form from a common gas cloud. As the resulting cluster disperses, its stars dissolve into the Galactic background yet maintain some common kinematical identity which may be used to identify members of a particular stellar group. Such moving groups represent a transition between bound clusters and field stars, and are probably chemically homogeneous (see \\citet{de-silva07} and references therein). Other moving groups appear to result from resonances in the galactic disk (eg \\citet{Dehnen98}), and others possibly as the debris of accreted and disrupted dwarf galaxies \\citep{Navarro04}.  We were interested to see whether the Kapteyn group members were chemically similar, because this would be a pointer to the group's origin.  It turned out that some of the group stars show chemical peculiarities similar to those seen in $\\omega$ Cen, and we will argue that the Kapteyn group may be part of the $\\omega$ Cen debris.  Concentrations of metal-poor stars at L$_z$ values near those of $\\omega$ Cen and the Kapteyn group have appeared in several recent studies.  \\citet{Dettbarn07} investigated substructure in a sample of stars with [Fe/H]$ < -1$ from the Beers et al (2000) catalog and identified a structure (their feature K) which appears to be related to the Kapteyn star group.  The study of 246 metal-poor stars with accurate kinematics by \\citet{Morrison09} also showed a substantial number of stars at weakly retrograde values of L$_z$, similar to that of $\\omega$ Cen. \\citet{Klement09} find evidence of Kapteyn stream stars in their study of halo streams from the SDSS DR7.  This paper presents a kinematic and chemical analysis of 17 members of the Kapteyn moving group.  Section \\ref{kins} presents a kinematic analysis of the group.  In Section \\ref{obs} the observations and reduction procedures are outlined, and Section \\ref{abunds} gives the details of the chemical abundance study.  The summary and conclusions are given in Section \\ref{sumandconc}. ",
        "conclusions": "\\label{sumandconc}  The kinematic and abundance analysis in this study suggests that at least our 14 members of the Kapteyn group and potentially many more stars in retrograde orbits which were not observed in this study, could be remnants of tidal debris stripped from the parent galaxy of $\\omega$ Cen, or even from the cluster itself,  during its merger with the Galaxy.  Our study provides the first detailed chemical evidence of field stars that appear to be both kinematically and chemically related to $\\omega$ Cen. It may lend weight to the view that $\\omega$ Cen is the remnant nucleus of a disrupted dwarf galaxy which was accreted by the Milky Way, by providing chemical evidence of tidal debris among the Galactic field stars.  The three-banded structure seen in the Lindblad diagrams of Figures 3 and described in Section 2.1 may be indicative of stars shed with different energies from different wraps of the decaying orbit of the parent galaxy around the Milky Way. The reader is referred to Figure 6 in Meza et al. (2005) which shows several distinct E--L$_z$ curves from their numerical simulations of merger debris.  In this study there is currently no evidence for an abundance difference between stars of different wraps of the orbit.  A study involving higher quality data and a larger number of stars from these wraps is underway. The presence of this banded structure, along with the present Galactic radius of $\\omega$ Cen, suggest the original parent galaxy was relatively massive in order for dynamical friction to have the required effect, and was also relatively dense in order to survive the Galactic tidal stresses in to the current orbital radius of $\\omega$ Cen.  What else can we infer about the parent galaxy?  If our stars are debris from the parent galaxy, then the galactic metallicity-luminosity relationship and the mean metallicity of the sample ($\\langle {\\rm [Fe/H]} \\rangle = -1.5$) indicate that the parent galaxy's luminosity would have been about $M_V \\sim -11$.  Its luminosity would have been about $2 \\times 10^6$ L$_\\odot$, and its stellar mass $\\sim 4 \\times 10^6$ M$_\\odot$. This stellar mass is comparable to the present stellar mass of $\\omega$ Cen itself.  The total mass of the parent galaxy could have been as high as $5 \\times 10^8$ M$_\\odot$ if it were a dwarf galaxy with a dark matter content similar to the Fornax dSph (e.g \\citet{Walker06}).  What are the chances of finding $\\omega$ Cen debris stars in the solar neighborhood from such a low-mass galaxy? We can use the Meza et al. simulation as a guide, although a detailed comparison is not appropriate as these authors note, because their end-product galaxy is not like the Milky Way. In summary, the debris of their satellite is well mixed azimuthally by redshift 0.48, at least within 10 kpc radius from the center of their parent galaxy, and is confined to a disk-like layer. The satellite is disrupted in its last three perigalactic passages, leaving a significant amount of substructure in E and L$_z$. We focus on the lowest $L_z$ substructure in their simulation, as a qualitative counterpart to the proposed $\\omega$ Cen debris. Most of our stars lie in a volume of radius $\\sim 500$ pc around the sun, and in a broad range of retrograde L$_z$. For comparison, we can estimate the fraction of the Meza et al. debris which would lie within our $500$ pc volume and within their lowest L$_z$ substructure: it is about $1.5 \\times 10^{-4}$. We could therefore expect to find about 700 debris stars in our volume and within the lowest angular momentum substructure, if we use the Meza et al.  simulations as a guide and assume that the typical star in this population has a mass of about 0.8 M$_\\odot$. In the Dinescu region shown in Figure 3, we have compiled a total of about 100 stars.  Even if all of them turn out to be $\\omega$ Cen debris stars, then this comparison may indicate that there are plenty more $\\omega$ Cen debris stars to be found nearby.  We should also consider the possibility that the Kapteyn group stars were stripped from the cluster itself as its outer regions were disrupted by the Galactic tidal field. This would be the immediate inference from the the chemical similarity which we have demonstrated between the Kapteyn group stars and  Cen stars. This seems unlikely, because the in-spiralling time under the influence of dynamical friction is in excess of the Hubble time for a cluster with the mass of $\\omega$ Cen; also, we note that the Kapteyn group stars are significantly more energetic than the cluster itself. It seems likely that the group stars came from the body of the parent dwarf galaxy, and therefore that the chemical peculiarities of $\\omega$ Cen are shared by its parent, and may well have originated in gas which was funneled into the cluster, possibly over an extended period, as suggested by \\citet{Norris97}.  Detailed numerical models of the chemical and dynamical evolution of $\\omega$ Cen as its galaxy loses mass in the Galactic tidal field are needed.  {\\bf Acknowlegements} ECW is funded by Australian Research Council grant DP0772283. KCF acknowledges his late colleagues Olin Eggen and Alex Rodgers for their work on the Kapteyn's star group which prompted this paper.  We would also like to add our thanks to the referee (D. Casetti-Dinescu) for many helpful comments and suggestions.. We thank Mike Bessell for his major improvements to the echelle spectrograph which made this work possible; Daniela Carollo for discussions which led to the association of the Kapteyn's star group with $\\omega$ Cen and for comments on an earlier draft of this paper; John Norris and Tim Beers for discussions and encouragement; and the SSO staff for maintaining the SSO instrumentation."
    },
    "0910/0910.3997_arXiv.txt": {
        "abstract": "Non-local thermodynamic equilibrium (NLTE) line formation for neutral and singly-ionized iron is considered through a range of stellar parameters characteristic of cool stars. A comprehensive model atom for Fe~I and Fe~II is presented. Our NLTE calculations support the earlier conclusions that the statistical equilibrium (SE) of Fe~I shows an underpopulation of Fe~I terms. However, the inclusion of the predicted high-excitation levels of Fe~I in our model atom leads to a substantial decrease in the departures from LTE. As a test and first application of the  Fe~I/II model atom, iron abundances are determined for the Sun and four selected stars with well determined stellar parameters and high-quality observed spectra. Within the error bars, lines of Fe~I and Fe~II give consistent abundances for the Sun and two metal-poor stars when inelastic collisions with hydrogen atoms are taken into account in the SE calculations. For the close-to-solar metallicity stars Procyon and $\\beta$~Vir, the difference (Fe~II - Fe~I) is about 0.1~dex independent of the line formation model, either NLTE or LTE. We evaluate the influence of departures from LTE on Fe abundance and surface gravity determination for cool stars. ",
        "introduction": "Iron plays an outstanding role in studies of cool stars thanks to quite numerous lines in the visible spectrum, which are easy to detect even in very metal-poor stars. Iron serves as a reference element for all astronomical research related to stellar nucleosynthesis and chemical evolution of the Galaxy. Iron lines are used to determine the surface gravity, $\\logg$, and the microturbulence $\\xi$ of stellar atmospheres. In the atmosphere with $\\Teff > 4500$~K, neutral iron is a minority species. The ionization equilibrium between Fe~I and Fe~II and the excitation equilibrium of Fe~I easily deviate from thermodynamic equilibrium. Since the beginning of the 1970s a number of studies attacked the problem of non-local thermodynamic equilibrium (NLTE) for Fe (e.g., \\cite[Athay \\& Lites (1972)]{Athay1972}, \\cite[Thevenin \\& Idiart (1999)]{Thevenin1999}, \\cite[Gehren \\etal\\ (2001)]{Gehren2001}). However, a consensus on the expected magnitude of the NLTE effects was not reached.  In this study, we update the model atom of Fe I-II treated by \\cite[Gehren \\etal\\ (2001)]{Gehren2001} (hereafter Paper~I) and apply it to analysis of the Fe spectrum in the Sun and selected cool stars with the aim of empirically constraining the role of inelastic collisions with hydrogen atoms in the SE of Fe~I-II.  \\firstsection ",
        "conclusions": "\\begin{itemize} \\item Completeness of model atom for Fe~I is important for a correct calculation of the Fe~I/Fe~II ionization equilibrium in the atmosphere of cool stars. \\item Thermalizing processes not involving electron collisions have to be included in the SE calculations for Fe~I-II. Collisions with hydrogen atoms could be good candidates for such processes. \\item Fe I is affected by significant NLTE effects for giants and very metal-poor stars. \\item Only minor departures from LTE are obtained for Fe II. \\end{itemize}  \\bigskip  {\\it Acknowledgements.} L.M. acknowledges a partial support from the International Astronomical Union, the Russian Foundation for Basic Research (08-02-92203-GFEN), and the Russian Federal Agency on Science and Innovation (02.740.11.0247) of the participation at the IAU XXVII General Assembly. This study is supported by the Deutsche Forschungsgemeinschaft (GE 490/34.1). A.K. acknowledges support by the Swedish Research Council (VR)."
    },
    "0910/0910.3445_arXiv.txt": {
        "abstract": "We present local extrema studies of two models that introduce a preferred direction into the observed cosmic microwave background (CMB) temperature field.  In particular, we make a frequentist comparison of the one- and two-point statistics for the dipole modulation and ACW models with data from the five-year \\textit{Wilkinson Microwave Anisotropy Probe} (\\WMAP). This analysis is motivated by previously revealed anomalies in the \\WMAP\\ data, and particularly the difference in the statistical nature of the temperature anisotropies when analysed in hemispherical partitions.  The analysis of the one-point statistics indicates that the previously determined hemispherical variance difficulties can be apparently overcome by a dipole modulation field, but new inconsistencies arise if the mean and the $\\ell$-dependence of the statistics are considered.  The two-point correlation functions of the local extrema, $\\xi_{\\rm TT}$ (the temperature pair product) and $\\xi_{\\rm PP}$ (point-point spatial pair-count), demonstrate that the impact of such a modulation is to over-`asymmetrise' the temperature field on smaller scales than the wave-length of the dipole or quadrupole, and this is disfavored by the observed data. The results from the ACW model predictions, however, are consistent with the standard isotropic hypothesis. The two-point analysis confirms that the impact of this type of violation of isotropy on the temperature extrema statistics is relatively weak.  From this work, we conclude that a model with more spatial structure than the dipole modulated or rotational-invariance breaking models are required to fully explain the observed large-scale anomalies in the \\WMAP\\ data. ",
        "introduction": "\\label{sec_intro}  Local extrema in the CMB have been extensively studied in the context of hotspots (peaks) and coldspots (troughs) arising in Gaussian random fields \\citep{bond_etal_1987, vittorio_etal_1987}. Additional statistics related to such local extrema, such as the Gaussian curvature and temperature-correlation function, have been investigated as a means to distinguish the geometry of the universe and test the Gaussian hypothesis for the nature of the initial conditions \\citep{barreiro_etal_1997, heavens_etal_1999, heavens_etal_2001}. Such statistical techniques have subsequently been applied to several datasets, and most notably by \\citet{kogut_etal_1995, kogut_etal_1996} with the $COBE$-DMR data. Observations from the \\textit{Wilkinson Microwave Anisotropy Probe} (\\WMAP) currently provide the most comprehensive, full-sky, high-resolution CMB measurements to-date, and studies on the local extrema properties thereof have been undertaken \\citep{LW04, LW05, tojeiro_etal_2006}.  More recently in \\citet{hou_etal_2009}, we extensively analyzed the statistical properties of both the one- and two-point statistics of local extrema in the five-year \\WMAP\\ temperature data. Such extrema are defined as those pixels whose temperature values are higher (maxima) or lower (minima) than all of the adjacent pixels \\citep{wandelt_etal_1998}. We considered only that part of the sky outside of the \\WMAP\\ KQ75 mask and its subsequent partition in either Galactic or Ecliptic coordinates into northern and southern hemispheres (hereafter GN, GS, EN, ES).  A frequentist comparison with the predictions of a Gaussian isotropic cosmological model that adopted the best-fit parameters from the \\WMAP\\ team was then made. The hypothesis test indicated a low-variance of both local maxima and minima in the Q-, V- and W-band data that was inconsistent with the Gaussian hypothesis at the 95\\% C.L. The two-point analysis showed that the observed temperature pair product at a given threshold $\\nu$, $\\xi_{\\rm TT}(\\theta,\\nu > 1,2)$, indicates a $3\\sigma$ level `suppression' on GN and EN, whereas $\\xi_{\\rm TT}(\\theta<20^{\\circ},\\nu < \\infty)$ is suppressed on the full-sky and both northern hemisphere partitions.  The latter is also the case for the point-point spatial pair-count function $\\xi_{\\rm PP}(\\theta<20^{\\circ},\\nu > 1,2)$. Intriguingly, the statistics showed an $\\ell$-dependence such that consistency with the Gaussian hypothesis was achieved once the first 5 or 10 best-fitting multipoles were subtracted, implying that the anomalies may be connected to features of the large-scale multipoles.  The local extrema anomalies therefore provide further evidence of a hemispherical asymmetry that was originally revealed using the power spectrum \\citep{hansen_etal_2004, eriksen_etal_2004, hansen_etal_2008}, the N-point correlation functions \\citep{eriksen_etal_2005} and Minkowski functionals \\citep{park_2004}. This can be interpreted as a violation of the cosmological principle of isotropy.  Theoretically, \\citet{gordon_etal_2005} proposed a mechanism of spontaneous isotropy breaking in which the long-wavelength modes of a mediating field couple non-linearly to the CMB perturbations. These fluctuations appear locally as a gradient and imprint a preferred direction on the sky. An implementation of a multiplicative modulation field of the intrinsic anisotropy with a single preferred direction was fitted to the three-year \\WMAP\\ observations in \\citet{eriksen_etal_2007} . The so-called \\lq dipole modulation field' (dmf) was detected at a significant confidence level, and confirmed at even higher significance in the five-year \\WMAP\\ data by \\citet{hoftuft_etal_2009}.  Another mechanism that violates rotational invariance was proposed by \\citet{acw_2007}. The ACW model picks out a preferred direction during cosmological inflation to modify the power spectrum of primordial perturbations, $P(k)$. Its imprint on the CMB can then be described by the covariance matrix of spherical harmonic coefficients, $\\langle a_{\\ell m}a^{*}_{\\ell'm'}\\rangle$. \\citet{groe_hke_2009} estimated the parameters of the ACW model in an extended CMB Gibbs sampling framework from the five-year \\WMAP\\ data. The posterior distribution of the parameters indicated a convincing detection of isotropy violation at $3.8\\sigma$ significance in the W-band data.  In this paper, we make a frequentist comparison of the dmf and ACW models with the five-year \\WMAP\\ maps of temperature anisotropy using a large number of Monte-Carlo simulations. In particular, we consider the one- and two-point statistics of local extrema and evaluate whether the \\WMAP\\ data are more consistent with these models as opposed to the standard Gaussian cosmological models against which various anomalies have been claimed. A rigorous hypothesis test methodology is applied to establish the significance at which the models impact the previous conclusions regarding \\WMAP\\ data and the violation of isotropy. In addition, the $\\xi_{\\rm TT}$ and $\\xi_{\\rm PP}$ correlation functions are utilised to further confirm the findings of the one-point analysis and to study the scale-dependence of the local extrema both in real and spherical harmonic space.  It might be considered that the statistical significance of the local extrema anomalies alone is insufficient to warrant an analysis of the anisotropic models, in particular given the additional parameters required to specify them. However, that they provide significantly better fits to the \\WMAP\\ data than the isotropic model has been demonstrated by independent Bayesian power spectrum analysis. It is clearly then quite legitimate to consider the local extrema statistics in terms of these improved models, and to determine whether the observed local extrema are also more consistent with such cosmological prescriptions.  In fact, we will find that there is little evidence that the extrema anomalies are remedied by these models, and indeed in one case we show that the anomalous behaviour becomes more significant.  This paper is organised as follows. In Section~\\ref{subsec_wmap_data}, we present an overview of the \\WMAP\\ data and the instrumental properties that must be involved in Gaussian simulations to enable an unbiased comparison with the real data. The two models and the algorithms for simulating them are briefly reviewed in Section~\\ref{subsec_2model}. Section~\\ref{subsec_analysis} prescribes the technique used for further data-processing and introduces the one- and two-point analysis methods. The results are reported in Section~\\ref{sec_results} before we present our conclusions in Section~\\ref{sec_conclusion}. ",
        "conclusions": "\\label{sec_conclusion}  In this paper, the statistical properties of local temperature extrema in the 5-year \\WMAP\\ data have been compared with two models that break rotational invariance -- dipole modulation and the ACW model. Such a comparison was motivated by the fact that our previous analysis of local extrema statistics \\citep{hou_etal_2009} demonstrated an unlikely hemispherical asymmetry in the GN and EN as a consequence of an extremely low-variance compared to the expectations of a Gaussian isotropic scenario. Moreover, the 5-year \\WMAP\\ data indicate significant detections of these models on the basis of power spectrum analyses.  We employ Gaussian MC simulations encoding the features of these two models to establish the statistical basis for testing whether the breaking of rotational invariance by the models can afford a satisfactory explanation of local extrema anomalies. As in previous work, both one and two-point statistics, as well as their dependence on large angular-scale modes have been studied. Both models are parameterised by a set of amplitudes and preferred directions as established by independent Bayesian analyses in \\citet{hoftuft_etal_2009} and \\citet{groe_hke_2009}. Our analysis has been carried out by sampling the former within their 95\\% confidence intervals, whilst imposing the fixed preferred direction found for each band, since the latter are estimated to be robust for each model. The V and W-band data are considered here.  In fact, the local extrema studies based on one and two-point statistics indicate that neither model provides a satisfactory solution to observed properties. In particular, the ACW model is probably the least interesting in this context -- both the one-point analysis and two-point studies show similar features to the isotropic results, including their $\\ell$-dependence. The dipolar modulation field, however, may itself be constrained by our analysis.  Results determined from division of the data into hemispheres implies that a dmf with significant amplitude may suppress the observed variance anomalies. In particular, the $\\lrmv=2$ values on EN and ES hemispheres indicates consistency of the data with the model. However, after subtraction of the first 10 multipoles, new problems arise in the southern hemispheres, where the observed variance of local extrema is now much lower than the model prediction. This is suggestive that the dmf may not have the correct profile over the full-sky in order to reconcile the variance properties of the observed local extrema with the Gaussian anisotropic simulations, and that the amplitude of the effect may also be a function of $\\ell$. Moreover, the observed mean statistics, which showed few anomalies when compared to the Gaussian isotropic model, are not consistent and contradict the variance results by requiring a weaker amplitude.  Our analysis indicates that neither a simple dipole modulated field model nor a rotational-invariance breaking model can satisfactorily explain the local extrema anomalies present in the \\WMAP\\ data. On the one side, the rotational-invariance model affects the large-scale temperature amplitudes too little to significantly affect the local extrema statistics. On the other side, a modulation type model needs a more elaborate spatial structure than a simple dipole to fully fit the data. These issues remain interesting for further investigation."
    },
    "0910/0910.5169_arXiv.txt": {
        "abstract": "We perform three-dimensional relativistic hydrodynamical simulations of the coalescence of strange stars and explore the possibility to decide on the strange matter hypothesis by means of gravitational-wave measurements. Self-binding of strange quark matter and the generally more compact stars yield features that clearly distinguish strange star from neutron star mergers, e.g. hampering tidal disruption during the plunge of quark stars. Furthermore, instead of forming dilute halo structures around the remnant as in the case of neutron star mergers, the coalescence of strange stars results in a differentially rotating hypermassive object with a sharp surface layer surrounded by a geometrically thin, clumpy high-density strange quark matter disk. We also investigate the importance of including nonzero temperature equations of state in neutron star and strange star merger simulations. In both cases we find a crucial sensitivity of the dynamics and outcome of the coalescence to thermal effects, e.g. the outer remnant structure and the delay time of the dense remnant core to black hole collapse depend on the inclusion of nonzero temperature effects. For comparing and classifying the gravitational-wave signals, we use a number of characteristic quantities like the maximum frequency during inspiral or the dominant frequency of oscillations of the postmerger remnant. In general, these frequencies are higher for strange star mergers. Only for particular choices of the equation of state the frequencies of neutron star and strange star mergers are similar. In such cases additional features of the gravitational-wave luminosity spectrum like the ratio of energy emitted during the inspiral phase to the energy radiated away in the postmerger stage may help to discriminate coalescence events of the different types. If such characteristic quantities could be extracted from gravitational-wave signals, for instance with the upcoming gravitational-wave detectors, a decision on the strange matter hypothesis and the existence of strange stars should be possible. ",
        "introduction": "The equation of state (EoS) of high-density matter and therefore the true nature of compact stars has been a mystery since the discovery of these objects. Besides the possibility of a neutron-proton composition more exotic phases have been proposed to appear in the cores of compact stars (see e.g. \\cite{2007ASSL..326.....H} for a review, for early considerations of quark stars see~\\cite{1969NCimL...2...13I,1970PThPh..44..291I}). The formulation of the strange matter hypothesis \\cite{PhysRevD.4.1601,PhysRevD.30.272} introduced the possibility that compact stars are so-called strange stars (SSs) \\cite{1986A&A...160..121H,1986ApJ...310..261A}. According to this hypothesis the absolute ground state of matter is formed by a quark phase consisting of up, down and strange quarks, and compact stars are self-bound objects composed of this kind of matter. It turns out that SSs are in many ways similar to neutron stars (NSs), e.g. in the mass range and compactness, and therefore they can be considered as an alternative explanation for compact stars \\cite{1986A&A...160..121H,1986ApJ...310..261A,1996csnp.book.....G,2007ASSL..326.....H,2005PrPNP..54..193W}. Throughout this article the term compact star refers to either a NS or a SS.  Theoretically as well as observationally, the question whether the strange matter hypothesis is true and whether (at least some) NSs are actually SSs, is an open issue. Besides the determination of compact star parameters, several possibilities have been proposed to explore the observational consequences of the strange matter hypothesis, including the possibility that small lumps of this strange quark matter (SQM) with baryon numbers of about $10^2$ and more might be abundant in the cosmic ray flux (see e.g. \\cite{2005PhRvD..71a4026M}). Experiments like AMS-02, which is planned to be installed on the International Space Station in 2010, have the potential to detect these so-called strangelets \\cite{ams,2004JPhG...30S..51S}. Also, SQM might be detected with the LHC at CERN \\cite{SchaffnerBielich:1996eh,Greiner:1987tg,Spieles:1996is}. For a review on these considerations and other observational implications of the strange matter hypothesis see \\cite{2005PrPNP..54..193W,2006JPhG...32S.251F}.  With the current and upcoming GW detectors like LIGO \\cite{:2007kva} and VIRGO \\cite{Acernese:2006bj} and the prospect of a detection of signals from compact object binaries \\cite{2004ApJ...601L.179K}, the question naturally arises if the signal of mergering NSs could be distinguished from SS mergers, especially because compact object binaries are considered to be among the most promising sources for these detectors.  The systematic investigation of the imprint of the EoS on the gravitational-wave (GW) signal is still in its infancy. Most studies consider simplified EoSs like polytropes and try to explore the chances to measure some general compact star properties like the stellar radius (e.g. \\cite{1996PhRvD..54.7261Z,2009arXiv0901.3258R}). It is not clear whether a decision on the strange matter hypothesis could be made on the basis of such measurements, because the compactness of SSs can be in the same range as that of NSs. First attempts of exploring the consequences of microphysical EoSs including nonzero temperature effects in a relativistic treatment, necessary for reliable results, have been made by \\cite{2007PhRvL..99l1102O}. Also fully relativistic studies have been conducted in which temperature corrections were approximated by an ideal gas component added to a zero-temperature microphysical EoS (e.g. \\cite{2006PhRvD..73f4027S} and preceding works).  Strange stars as sources of GWs have been considered in a rotational equilibrium approach investigating the (final phase) of the inspiral of SS binaries \\cite{PhysRevD.71.064012,GondekRosinska:2008nf}. Quasiequilibrium orbits were constructed to determine the innermost stable circular orbit (ISCO). It was found that the frequency of the GWs at the ISCO depends on the compactness of the SSs. The orbital evolution and the associated GW emission of SS-black hole binaries has been estimated within a semianalytic model in \\cite{2004JPhG...30.1279P}. Moreover, the signals from various instabilities of rotating SSs have been worked out in \\cite{2000PhRvL..85...10M,2002MNRAS.337.1224A,GondekRosinska:2003iy} (see also references therein). In \\cite{2008PhRvL.101r1102F} it was found that the frequencies of the g-mode in newly born SSs are significantly lower than those of NSs. Differences of the fundamental pressure modes are discussed in \\cite{2007GReGr..39.1323B} where a discrimination was found only to be possible if additional information like the mass was available. Nevertheless it is not clear whether these discriminating features are relevant as the corresponding GW signals might be too weak for measurements.  In this article we examine the GW signals emitted by merging SSs. We explore two different EoSs representing different possibilities of SQM properties and consider the inspiral phase, the final plunge, and the postmerger stage. We compare the GW characteristics of SS coalescence with those of ordinary NSs and discuss if GW observations could be decisive on the strange matter hypothesis. For that we conduct three-dimensional relativistic hydrodynamical simulations that allow us to follow the evolution of the whole merging process and to extract observational signatures. To our knowledge this is the first study considering these events in dynamical simulations; only results for mergers of SS-black hole systems have been so far reported in \\cite{2002MNRAS.335L..29K}, where Newtonian hydrodynamics were used and the black hole (BH) was modeled by a pseudorelativistic potential \\cite{1980A&A....88...23P}.  The paper is structured as follows: Sect.~\\ref{sec:num} will introduce the underlying model and numerical methods. In Sect.~\\ref{sec:eos} the properties of the EoSs used in this study will be described. The simulations and some general aspects will be discussed in Sect.~\\ref{sec:sim}. In Sect.~\\ref{sec:GW} the GW signals will be presented and in Sect.~\\ref{sec:con} we will close with a summary and our conclusions. ",
        "conclusions": "\\label{sec:con}  We have reported on simulations of SS mergers and discussed our results in comparison to NS coalescence events. In particular, we focussed on the GW signals and their potential to provide information about the existence of SSs.  We found that the dynamical behavior of SS mergers is fundamentally different, which can be understood by the higher compactness of SSs and their merger remnants, the self-binding of SQM and the different influence of thermal effects. While NS mergers form a dilute halo or toruslike structure around a dense, hypermassive, differentially rotating remnant, the remnant of merging SSs is bounded by a sharp surface as were the initial stars. Only by the formation of thin tidal arms relatively late in the evolution of the remnant, matter gets shed off the central object and forms a fragmented thin disk in the equatorial plane.  In order to estimate the importance of thermal effects during NS and SS mergers we compared our models with simulations where zero temperature was imposed by extracting the thermal energy as most efficient cooling would do. For both kinds of stars we observed a sensitive dependence of the dynamics of the system on thermal effects, affecting for instance the development and structure of the outer remnant parts of NS mergers. While the estimated relic torus mass after collapse of the remnant to a BH is in general higher if thermal effects were neglected, a nonzero temperature of NS matter leads to an inflated, dilute halolike torus in contrast to a much thinner, more disklike structure for $T=0$. For SS mergers we found a shorter time delay for the BH formation, whereas in the case of NSs this time interval is stretched as reported in \\cite{2008PhRvD..78h4033B}. This difference can be understood as a consequence of the additional gravitating effect of the thermal energy, which in the case of very dense SSs is not overcompensated by thermal pressure effects (different from the NS case).  The analysis of the GW signals emitted by SS mergers revealed that already by means of relatively simple characteristic features of the signal it may be possible to decide whether SS or NS mergers produced the emission. In particular, we found that the maximal frequency during the inspiral and the frequency of the ringdown of the postmerging remnant are in general higher in the case of SS mergers in comparison to NS mergers. Whether this criterion can be used to finally decide on the strange matter hypothesis depends on the particular stellar properties associated with the EoS of high-density matter. For a similar mass-radius relation within a certain mass range, meaning relatively compact NSs or less compact SSs (the LS-MIT60 scenario), the determination of these frequencies might not be decisive. However, taking additional characteristics of the GW luminosity into consideration will allow for a discrimination. In this context we discussed the occurrence of a prompt collapse, the ratios of the GW energy emitted in the postmerger phase to the energy radiated away during the inspiral phase, the growth rate of the energy emission associated with the postmerger ringdown signal, and the appearance of a gap in the GW luminosity spectrum. In addition, it was shown in \\cite{2008arXiv0812.4248B} that cosmic ray experiments will yield information about the strange matter hypothesis and may clarify the situation, especially in the case when GW signals are not conclusive as in the LS-MIT60 scenario. Therefore we expect that despite the rather uncertain event rate the upcoming advanced GW detectors LIGO and VIRGO may provide valuable data to decide about this long-standing question, and it appears likely that one will thus gain fundamental information about the properties of high-density matter. Moreover, future GW detectors like the Einstein telescope \\cite{et} and the DUAL detector \\cite{dual} will have a higher sensitivity in the high-frequency domain above 1~kHz, where characteristic features occur. Signals measured with these future instruments will therefore allow for a more detailed analysis.  Finally, we would like to point out that also other forms of self-bound matter have been discussed in the literature (see e.g.~\\cite{1977JETPL..25..465V,1990PhR...192..179M,2007ASSL..326.....H}). For instance pion condensates may lead to stellar objects similar to SSs, and also abnormal nuclei similar to strangelets may exist in this case. Therefore, the study presented in this work should be generalized to these states of matter, where quantitative considerations require the knowledge of the specific EoS. However, for stellar properties comparable to those of SSs we expect qualitatively the same behavior and the same GW features as for SSs. From this discussion it is clear that only a multimessenger approach can bring a final answer. In addition to GW measurements and cosmic ray experiments, this may in particular include the observation of isolated compact stars to derive the cooling history of these objects, which is significantly different for NSs, SSs and self-bound pion-condensed stars (see~\\cite{1999paln.conf.....W,2000ApJ...533..406B,2004ARA&A..42..169Y,2005PhRvC..71d5801G}).  Several improvements and extensions may supplement this first study of SS mergers in the future. Besides refinements in the methodical approach like a fully relativistic treatment and the inclusion of magnetic fields (only important during the merging phase if very strong, but must be expected to grow afterwards and affect the ringdown phase) and radiative effects, one would also like to account for the effects of quark interaction and color superconductivity, employ models beyond the MIT bag model, and discuss other forms of self-bound matter. Moreover, a determination of the modes excited in the merger remnants may be a promising way to develop a better understanding of the origin of the characteristic properties of the GW spectra. Furthermore a study of SS-BH mergers will yield further insights about the behavior of such systems in contrast to NS-BH binaries."
    },
    "0910/0910.2244.txt": {
        "abstract": "We discovered a sample of 250 Ly$\\alpha$ emitting (LAE) galaxies at $z\\simeq2.1$ in an ultra-deep 3727{\\AA} narrow-band MUSYC image of the Extended Chandra Deep Field-South.  LAEs were selected to have rest-frame equivalent widths (EW) $>20$~{\\AA} and emission line fluxes $F_{Ly\\alpha} >2.0 \\times 10^{-17}$~erg~ cm$^{-2}$~s$^{-1}$, after carefully subtracting the continuum contributions from narrow-band photometry. The median emission line flux of our sample is $F_{Ly\\alpha} = 4.2 \\times10^{-17}$~erg~cm$^{-2}$~s$^{-1}$, corresponding to a median Ly$\\alpha$ luminosity L$_{Ly\\alpha} = 1.3 \\times10^{42}$~erg~s$^{-1}$ at $z\\simeq2.1$. At this flux our sample is $\\geq90$~\\% complete. Approximately 4\\% of the original NB-selected candidates were detected in X-rays by Chandra, and 7\\% were detected in the rest-frame far-UV by GALEX; these objects were eliminated to minimize contamination by AGN and low-redshift galaxies. At L$_{Ly\\alpha} \\geq 1.3 \\times10^{42}$~erg~s$^{-1}$, the equivalent width distribution is unbiased and is represented by an exponential with scale-length 83~$\\pm10$~{\\AA}.  Above this same luminosity threshold, we find a number density of $1.5\\pm0.5\\times 10^{-3}$~Mpc$^{-3}$. Neither the number density of LAEs nor the scale-length of their EW distribution show significant evolution from $z\\simeq3$ to $z\\simeq2$. We used the rest-frame UV luminosity to estimate a median star formation rate of 4 M$_{\\odot}$~yr$^{-1}$. The median rest-frame UV slope, parametrized by the color $B-R$, is that typical of dust-free, 0.5-1 Gyr old or moderately dusty, 300-500 Myr old population. Approximately 30\\% of our sample is consistent with being very young (age~$<100$ Myr) galaxies without dust. Approximately 40\\% of the sample occupies the $z\\sim2$ star-forming galaxy locus in the $UVR$ two color diagram, but the true percentage could be significantly higher taking into account photometric errors. Clustering analysis reveals that LAEs at $z\\simeq2.1$ have $r_0=4.8\\pm0.9$ Mpc, corresponding to a bias factor $b=1.8\\pm0.3$. This implies that $z\\simeq2.1$ LAEs reside in dark matter halos with median masses log(M/M$_{\\odot})= 11.5^{+0.4}_{-0.5}$, which are among of the lowest-mass halos yet probed at this redshift. We used the Sheth \\& Tormen conditional mass function to study the descendants of these LAEs and found that their typical present-day descendants are local galaxies with L$^*$ properties, like the Milky Way. ",
        "introduction": "The search for high-redshift star-forming galaxies advanced rapidly with the introduction of the Lyman Break Galaxies (LBG) technique (\\citealt{Gu:1990},  \\citealt{StHm:1992}, \\citealt{Steidel:1999}) that takes advantage of the lack of flux at wavelengths shorter than the Lyman break at 912~{\\AA} ~rest frame due to absorption of ionizing photons by neutral hydrogen, located in stellar atmospheres, in the interstellar medium (ISM), and in the intergalactic medium (IGM) between galaxies. At $z=3$ the break is located in the observed U band and at higher redshift it moves into optical and infrared bands. This has allowed an exploration of star-forming galaxies at redshifts $3\\leq z\\leq8$ via imaging from ground and space (e.g. \\citealt{Steidel:2003}, \\citealt{Bouwens:2006}, \\citealt{Ouchi:2008}) and spectroscopy on 8--10 meter telescopes (e.g. \\citealt{Shapley:2001},  \\citealt{Shapley:2003}). A significant fraction of high redshift LBGs show the Ly$\\alpha$ line in emission (\\citealt{Shapley:2001}). This emission offers additional information about the process of star formation inside these galaxies and radiative transfer in their ISM.  Looking for galaxies with Ly$\\alpha$ in emission has become an important photometric technique that permits us to find faint (R $\\sim$ 27) star-forming galaxies at high redshift. This technique consists of comparing the flux density measured in a narrow-band filter, revealing observed-frame Ly$\\alpha$ emission to that found in the broad-band filters, representing the continuum. Thanks to the intensity of this emission line, the resulting Ly$\\alpha$ emitting (LAE) galaxies provide a special population of high redshift galaxies. The properties of LAEs have been extensively studied at $z\\ge3$ (e.g. \\citealt{Ouchi:2005}, Venemans et al. 2005, \\citealt{Gawiser:2006b}, \\citealt{Gronwall:2007}, ~\\citealt{Nilsson:2007}).  LAE samples are composed primarily of galaxies fainter in the continuum than LBGs; Ly$\\alpha$ Emitting galaxies therefore probe the lowest bolometric luminosities at high redshift.  Theoretical models, that include radiative transfer inside star-forming galaxies (\\citealt{V:2006}, \\citealt{SV:2008}, \\citealt{V:2008},  \\citealt{Atek:2009}), were also developed to understand how Ly$\\alpha$ photons form in HII regions and then escape the galaxy, depending on resonant scattering by neutral hydrogen, dust absorption and velocity dispersion in the interstellar medium. The amount of dust and the interstellar medium geometry can affect the escape of Ly$\\alpha$ photons and hence the shape of the line. Clumpy media could permit Ly$\\alpha$ photons to escape, even if the galaxy is not dust-free (\\citealt{Neufeld:1991}, \\citealt{Fin:2008}, \\citealt{Fyougold:2009}).  Spectral Energy Distribution (SED) fitting of the stacked multi-wavelength photometry of $z\\simeq3$ LAEs (\\citealt{Gawiser:2007}, \\citealt{Lai:2008}) shows they are a young (median starburst age of $\\sim20$ Myr), low stellar mass (M$\\sim10^9$ M$_{\\odot}$), modest SFR (median SFR $\\sim2$ M$_{\\odot}$~yr$^{-1}$), low dust (A$_V\\leq$0.2) population of galaxies in an active phase of star formation. SEDs have also shown older population best fits for subsamples of LAEs at $z>3$ (\\citealt{Pirzkal:2007}, \\citealt{Ono2009}, \\citealt{Nilsson:2009}). SED fitting of individual galaxies showed older ages, higher stellar mass and more dust for continuum-bright LAEs drawn from LBG samples (\\citealt{Shapley:2001}, \\citealt{Tapken:2007}, \\citealt{Pentericci:2009}). \\citet{Stiavelli2001} had also shown redder colors for some LAEs at $z\\simeq2.4$. Recently \\citealt{Nilsson:2009} presented the first results of observations of LAEs at $z\\simeq2.3$, inferring evolution in the properties from $z\\sim3$ to $z\\sim2$, with more diversity in photometric properties at $z\\simeq2.3$.  Clustering analysis of LAEs showed $z\\ge4$ LAE bias factors (\\citealt{kovac:2007}, \\citealt{Ouchi:2004}) expected for progenitors of massive elliptical galaxies in the local Universe, while $z\\simeq3.1$ LAEs could be progenitors of L$^*$ galaxies (\\citealt{Gawiser:2007}). Semi-analytical simulations were also able to reproduce these results (\\citealt{Orsi2008}).  Lower redshift observations, including clustering, will reveal evolution from high to low redshift. For this reason we were motivated to study LAE samples at redshift around 2. This will trace the star formation properties of this type of galaxy at the epoch of the peak of cosmic star formation density (\\citealt{Madau:1998}, \\citealt{Giav2004}). It also promises to reveal $z\\sim0$ descendants of LAEs at $z\\simeq2.1$.  In this paper we describe the results from ultra-deep 3727~{\\AA} narrow-band MUSYC (MUlti-walengthSurvey Yale Chile, \\citealt{Gawiser:2006a}) imaging of the 998 arcmin$^2$ Extended Chandra Deep Field-South. In sections 2 and 3 we present the observations and the data reduction. In Section 4 we summarize the selection of the LAE sample and estimate the possible contaminants. We present the properties of the LAE sample in Section 5: number density, star formation rate, colors and clustering. In Section 6 we discuss the results and derive conclusions.  We assume a $\\Lambda$CDM cosmology consistent with WMAP 5-year results (\\citealt{Dunkley:2009}, their table 2), adopting the mean parameters  $\\Omega_m=0.26$, $\\Omega_{\\Lambda}=0.74$, H$_0=70$ km~sec$^{-1}$~Mpc$^{-1}$, $\\sigma_8=0.8$. ",
        "conclusions": "\\label{sec:discussion}  We imaged the ECDF-S using a NB3727 narrow-band filter, corresponding to the wavelength of Ly$\\alpha$ emission at $z\\simeq2.1$. Following the formalism described in the Appendix, we applied the color cut $UB$$-$NB3727$ > 0.73$ and additional significance criteria that yielded a sample of 250 LAEs. In our observation we achieve the same 5$\\sigma$ detection limit in Ly$\\alpha$ luminosity (log(L(Ly$\\alpha$))=41.8) as the sample of LAEs at $z\\simeq3.1$ (Gronwall et al. 2007, Gawiser et al. 2007). Therefore we are able to look for indications of evolution between $z\\sim$2 - 3. Concentrating on $z\\sim2$, we compare LAEs with  star-forming galaxies (Steidel's BX sample), which can also show the Ly$\\alpha$ line in emission. In many cases our analysis concentrates on the typical properties of the LAE sample as a whole; it is important to remember that there will always be cases of individual LAEs whose physical properties differ considerably from those of the typical LAE.  The magnitude distribution of LAEs at $z\\simeq2.1$ (Fig.~\\ref{fig:mag}) is consistent with that predicted by the $z\\simeq3.1$ LAE Ly$\\alpha$ luminosity function, but with about twice the normalization, i.e. total number density. As reported in \\S 5.1 we calculated a LAE number density of $3.1\\pm0.9\\times 10^{-3}$ Mpc$^{-3}$, taking into account the estimated incompleteness of the sample, an evolution in the number density of a factor of $2.1\\pm0.7$ versus $1.5\\pm0.3 \\times 10^{-3}$ Mpc$^{-3}$ reported by \\citeauthor{Gawiser:2007} (2007) at $z\\simeq3.1$. Our number density is consistent with the value, found by Nilsson et al. (2009) at $z\\simeq2.3$ when we restricted our analysis to objects matching their $\\sim2 \\times$ brighter luminosity limit. At the Ly$\\alpha$ luminosity limit, at which the sample is complete, we calculate a number density of $1.5\\pm0.5\\times 10^{-3}$ Mpc$^{-3}$, that implies an increasing factor of $1.4\\pm0.5$, consistent with that calculated for all the sample.  We derive the equivalent width distribution (\\S5.2), representative of the $z\\simeq2.1$ LAE sample in Fig.~\\ref{fig:ew}. As we can see in Fig.~\\ref{fig:ew}a, for log(L(Ly$\\alpha))\\geq42.1$ the sample is unbiased in the sense of rest-frame equivalent width versus Ly$\\alpha$ luminosity. We consider the unbiased brighter half of the sample to build the histogram in Fig.~\\ref{fig:ew}b. Fitting this distribution with an exponential law, this is consistent with that from Gronwall et al. (2007) for the sample at $z\\simeq3.1$ and broader than that found at $z\\simeq2.3$ by Nilsson et al. (2009). In Fig.~\\ref{fig:ew} we associated the value EW~$=400$~ {\\AA} to the objects characterized by an unphysical equivalent width (Appendix, equation \\ref{colorequation}). The objects with $EW_{rest-frame}>250$ present $UB>27$. Most of the objects in the sample with $EW_{rest-frame}<50$ also have log(L(Ly$\\alpha))<42.1$, meaning that their continuum flux boosted them above the narrow-band catalog detection limit. This behavior was less prevalent at $z\\simeq3.1$ by Gronwall et al. (2007), although the 5$\\sigma$ detection limit creates a similar trend, as shown by the blue curve in Fig.~\\ref{fig:ew}a. As it is described in the Appendix, we estimate the observed EW from the observed color $UB$$-$NB3727. Those estimations are in perfect agreement with those obtained from continuum flux density and Ly$\\alpha$ emission line flux. As described in Dayal et al. (2009), the measured EW at the border of the galaxies can be increased by the cooling of collisionally interstellar medium excited HI atoms, while the continuum almost remains unchanged, but intergalactic medium absorption can attenuate Ly$\\alpha$ flux and so decrease the observed EW.  The Ly$\\alpha$ luminosity reveals star formation activity inside a galaxy (\\S5.3).  Log(L(Ly$\\alpha))=42.1$, the median Ly$\\alpha$ luminosity of our sample, corresponds to SFR(Ly$\\alpha$) =  1.2 M$_{\\odot}$~yr$^{-1}$, as indicated by the dashed-dotted line of Fig.~\\ref{fig:sfr}. In the same figure we observe the range of SFR(UV) values. The median LAE at $z\\simeq2.1$ has a moderate SFR(UV) of $\\sim$ 4 M$_{\\odot}$~yr$^{-1}$. The ratio of $\\sim$1.5 in the values of SFR(UV)/SFR(Ly$\\alpha$), for the unbiased half of the sample, is caused by potentially complex radiative transfer of Ly$\\alpha$ photons in the dusty, possibly clumpy interstellar medium inside the galaxies \\citep{Atek:2009}. Given the overlap in clustering bias it is worth considering whether $z\\simeq2.1$ LAEs could populate the low (stellar) mass tail of continuum-selected star-forming galaxies at  $z\\sim2$. We find that the LAE SFR(UV) is 10 times lower than that calculated from UV continuum and H$\\alpha$ line emission by \\citeauthor{Steidel:2004} (2004) for star-forming galaxies at $z\\sim2$. The Kennicutt estimator, used to derive the star formation rate from UV continuum, assumes that the galaxy is at least 10$^7$ yr old with  roughly constant SFR.  We find (\\S5.4) that 240/250 (96\\%) of $z\\simeq2.1$ LAEs are blue ($B-R$)$<1$, with 73/250 (30\\%) having ($B-R$)$<0$.  This is in good agreement with the $z\\simeq3.1$ sample in both criteria. In fact at $z\\simeq3.1$, LAEs with $R<25$ have median color $B-R=0.53$ (\\citealt{Gronwall:2007}). Our result agrees with the findings of \\citet{Nilsson:2009} at $z\\simeq2.3$ in the fraction of LAEs having ($B-R$)$>0$, but their conclusion that most LAEs are ``red'' depended on considering all objects with rising spectra in $f_{\\nu}$ to be red. A reasonable split of galaxies into blue and red is achieved by using ($B-R$)$=1$ as the dividing line, and we suspect that the sample of \\citet{Nilsson:2009} will show similar properties when this is applied. In fact looking at their Fig. 4 and deriving the behavior of the color $B-R$ from the slope $\\beta(B-R)$, we see that their galaxies are essentially blue, based on our definition.  The appearance of bimodality in the LAE rest-UV color at $R<25$ is intriguing. The blue branch is presumably dominated by young, dust-free star-forming galaxies, since unobscured (blue) AGN should have been eliminated from our sample due to their X-ray emission.  The red branch may contain obscured (dust-reddened) AGN, galaxies with Ly$\\alpha$ emission from recent starbursts but an overall older or dustier stellar population and low-redshift interlopers that will be identified via follow-up spectroscopy. We calculated the evolutionary tracks of galaxies at z$\\sim$2 in the $U-B$ vs $B-R$ plane, generated using the GALAXEV (Bruzual \\& Charlot 2003) code for a constant star formation rate and a range of masses from 25 Myr to 1 Gyr,  parameters consistent with LAE SED fits.  We see that a 500 Myr old galaxy with dust absorption A$_V=0$ has color $B-R=0$. If it is star forming, the $U-B$ color, corrected for IGM absorption, is also close to zero. Increasing the age the color $B-R$ becomes slightly bigger than 0. However increasing the amount of dust, for example to A$_V\\sim$1, typical for reddened LBG, the star-forming galaxy can assume $B-R$=0.5-0.6. The color $B-R$=1 is achieved by galaxies with significantly more dust than that measured for typical star-forming populations. There is a smaller difference in $B-R$ between young ($<5 \\times10^8$ yr) and old ($>5 \\times10^8$ yr) star-forming populations than the difference produced by the increasing reddening. The observed median($B-R$)~=~0.16 is typical of star-forming galaxies with A$_V$~=~0 and ages of 0.5-1 Gyr or can be consistent with moderate A$_V$ and age 300-500 Myr. Approximately 30 \\% of our sample with negative $B-R$ color is consistent with being very young (age$<$100 Myr) galaxies without dust.  We divide in bins of 0.5 magnitude in $R$ and construct Table \\ref{tab:Rchar}, which shows the magnitude range, the median color, EW, SFR from Ly$\\alpha$ and SFR from the UV continuum. These values are transformed into intrinsic ones, taking into account the dust and gas amount (parametrized by stellar E(B-V) and E$_g$(B-V) ) and radiative transfer effects. The median colors lie inside the ``BX\" region except for the faintest bins which have large photometric uncertainties.  As expected we observe that the EW values are bigger for the objects that are fainter in the continuum. We calculated $EW_{rest-frame}>250$ for objects with $UB>27$. Statistical fluctuations related to such a faint continua can produce an over-estimation of the equivalent width of these objects. We observe that bright-continuum objects ($UB<24.5$) are also bright in Ly$\\alpha$ luminosity. For low-EW LAEs ($UB-NB3727$ just $\\simeq 0.73$), as expected, the SFR(UV) is significantly bigger than the SFR(Ly$\\alpha$). In the table we also report the standard deviations in the $R$ magnitude bins. In the last column the scatter error is less meaningful, because of the proportionality between $R$ flux density and SFR(UV). It is seen that the scatter is as big as the corresponding quantity. In $B-R$ color it is consistent with that was observed in Fig.~\\ref{fig:BR}.  The clustering analysis (\\S5.5) gives information about the LAEs at $z\\simeq2.1$ as a population and their evolution to redshift zero. In Fig.~\\ref{fig:cluster2} we see that LAEs at very high redshift ($z>4$, Ou sign, H09 sign) can evolve into massive LBG at $z\\sim$3 and also reach, in the local Universe, the bias factor typical of elliptical massive galaxies, corresponding to luminosity between 2.5 and 6.0 L* (as indicated by the points in the figure) and halo masses greater than $4.47 \\times 10^{13} $~M$_{\\odot}$.  Looking at $z\\sim3$ \\citep{Gawiser:2007}  LAEs were observed to be blue galaxies and to be characterized by lower clustering than other galaxy samples at that redshift. They can evolve into star-forming galaxies at $z\\sim2$ (A0 sign) and then to L$^*$ galaxies in the local Universe. At  $z\\simeq2.1$ we calculate a bias factor b~$=1.8\\pm$0.3 for our sample of LAEs. This value is consistent with that found using the conditional mass function for progenitors of  L$^*$ galaxies in the local Universe. It is also consistent with the value calculated for the subset of ``BX\" galaxies dimmest in $K$-band ($K_{Vega}>21.5$, \\citealt{Adelbergerb:2005}); that is low mass galaxies. This clustering result matches that of dark matter halos with median masses of log(M/M$_{\\odot})= 11.5^{+0.4}_{-0.5} $, which are some of the lowest halo masses  probed at this redshift. Our result shows that $z\\sim2$ LAEs could also be descendants of $z\\simeq3.1$ LAEs, depending on how long dust-free star formation occurs and on possible cyclical repetitions of star formation phases. As LAEs at $z\\simeq2.1$ are consistent with being the progenitors of present-day and L$^*$ galaxies at $z=0$, they are likely building blocks of local galaxies with properties similar to the Milky Way  and median halo mass  $\\geq 2\\times 10^{12}$ M$_{\\odot}$."
    },
    "0910/0910.0266_arXiv.txt": {
        "abstract": "Much of our understanding of dark matter halos comes from the assumption that the mass-to-light ratio ($\\Upsilon$) of spiral disks is constant.  The best way to test this hypothesis is to measure the disk surface mass density directly via the kinematics of old disk stars.  To this end, we have used planetary nebulae (PNe) as test particles and have measured the vertical velocity dispersion ($\\sigma_z$) throughout the disks of five nearby, low-inclination spiral galaxies:  IC~342, M74 (NGC~628), M83 (NGC~5236), M94 (NGC~4736), and M101 (NGC~5457).  By using H{\\sc i} to map galactic rotation and the epicyclic approximation to extract $\\sigma_z$ from the line-of-sight dispersion, we find that, with the lone exception of M101, our disks do have a constant $\\Upsilon$ out to $\\sim$3~optical scale lengths ($h_R$).  However, once outside this radius, $\\sigma_z$ stops declining and becomes flat with radius.   Possible explanations for this behavior include an increase in the disk mass-to-light ratio, an increase in the importance of the thick disk, and heating of the thin disk by halo substructure.   We also find that the disks of early type spirals have higher values of $\\Upsilon$ and are closer to maximal than the disks of later-type spirals, and that the unseen inner halos of these systems are better fit by pseudo-isothermal laws than by NFW models. ",
        "introduction": "Traditionally, spiral galaxies have three main components: a dynamically hot bulge, a dynamically cold disk (which can sometimes be divided into ``thin'' and ``thick'' components), and a mysterious dark matter halo.  However, measuring the mass of each component separately is a challenge, especially in a galaxy's extreme outer regions where the dark matter halo dominates.  Even in late-type spirals with minimal bulges, the mass profiles of the dark halos cannot be decoupled from the visible disk mass using rotation curves alone \\citep{bsk04}.  Most analyses have therefore relied on the ``maximal disk'' hypothesis, wherein the disk mass-to-light ratio ($\\Upsilon$) is assumed to be constant with radius, and as large as possible to account for the rotation in the inner parts of the galaxy \\citep[\\eg][]{kent86, pw00, sofue03}.  Is the disk $\\Upsilon$ really constant and maximal? Perhaps the best way to break the disk-halo degeneracy is to measure the surface mass of a disk directly from the $z$~motions of stars.  In systems where the primary stellar motion is circular, the phase-space distribution of stars in the direction perpendicular to the plane is fixed by the disk's gravitational potential \\citep{bt87}.  For example, in the Milky Way, star counts and velocity measurements toward the Galactic pole initially suggested that the stellar disk is roughly isothermal \\citep[\\eg][]{oort}.  If this were indeed the case, then the integrated face-on stellar velocity dispersion, $\\sigma_z$, would be related to the disk surface mass, $\\Sigma$, by \\begin{equation} \\sigma_z^2(R) = K G \\, \\Sigma(R) h_z, \\label{isothermal} \\end{equation} where $K = 2 \\pi$ is a constant, $G$ is the gravitational constant, and  $h_z$ is the scale height of the stars \\citep{vdK88}.   Since surveys of edge-on galaxies \\citep[\\eg][]{vdKs82, bm02} imply that $h_z$ is roughly constant with radius, measurements of $\\sigma_z$ can produce estimates of disk mass that are independent of a galaxy's rotation curve.  Absorption line studies have shown that, in the inner regions of spirals, $\\sigma_z$ generally follows the exponential law expected from a constant $\\Upsilon$ disk \\citep{b93, g+97, g+00}.  However, these surveys were limited by surface brightness, and do not extend more than $\\sim$2~scale lengths from the nucleus.  To probe the outer regions of disks, where the dark matter halo is more important, some other technique is required.  Planetary nebulae (PNe) are ideal tools for this purpose: PNe are extremely bright in [\\ion{O}{3}], abundant out to $\\gtrsim$5~disk scale lengths, measurable to $\\sim$2~\\kms\\ with fiber-fed spectrographs, and present in all stellar populations with ages between $10^8$ and $10^{10}$~yr. Moreover, these old disk stars are easy to distinguish from other emission-line sources via their distinctive [\\ion{O}{3}]-H$\\alpha$ ratio \\citep{p12}.  \\begin{deluxetable*}{cccccccccc} \\tabletypesize{\\scriptsize} \\tablecaption{Program Galaxies\\label{tabBasic}} \\tablewidth{0pt} \\tablehead{&\\colhead{Hubble}&&\\colhead{Distance}&\\colhead{$h_R$}&\\colhead{Disk}&&\\colhead{$v_{max}$}&\\colhead{Number of}&\\colhead{Survey} \\\\ \\colhead{Galaxy} &\\colhead{Type\\tablenotemark{a}} &\\colhead{$i$\\tablenotemark{b}} &\\colhead{(Mpc)\\tablenotemark{c}} &\\colhead{(kpc)\\tablenotemark{d}} &\\colhead{$\\mu_0$\\tablenotemark{e}} &\\colhead{$E(B-V)$\\tablenotemark{f}} &\\colhead{(\\kms)\\tablenotemark{g}} &\\colhead{PN Velocities} &\\colhead{Region} } \\startdata IC~342 &Scd  &$31^\\circ$  &$3.5 \\pm 0.3$ &4.24 &19.60 &0.558 &200 &99 &$4\\farcm 8$  \\\\ M74    &Sc   &$6.5^\\circ$ &$8.6 \\pm 0.3$ &3.17 &20.20 &0.070 &170 &102 &$4\\farcm 8$ \\\\ M83    &SBc  &$24^\\circ$  &$4.8 \\pm 0.1$ &2.45 &19.07 &0.066 &255 &162 &$18\\arcmin$ \\\\ M94    &Sab  &$41^\\circ$  &$4.4^{+0.1}_{-0.2}$ &1.22  &18.88 &0.018 &200 &127 &$5\\farcm 8$ \\\\ M101   &Scd  &$17^\\circ$   &$7.3 \\pm 0.5$ &4.99 &20.29 &0.009 &250 &60 &$8\\arcmin$ \\enddata \\tablenotetext{a}{From RC3} \\tablenotetext{b}{IC 342: \\citet{cth00}; M74: \\citet{fb+07}; M83: \\citet{p+01}; M94: \\citet{dB+08}; M101: \\citet{ZEH90}} \\tablenotetext{c}{From Paper~I, except for M101, which is from \\citet{M101PNe}; note that this distance has been updated to be consistent with $M^* = -4.47$ \\citep{p12}} \\tablenotetext{d}{IC~342: \\citet[][$I$-band]{wb03}; M74: \\citet[][$R$-band]{m04}; M83: \\citet[][$R$-band]{k+00}; M94: \\citet[][$3.6\\mu$m]{dB+08} and \\citet[][$R$-band]{g+09}; M101: \\citet[][$R$-band]{khc06}} \\tablenotetext{e}{In mag~arcsec$^{-2}$, corrected for galactic inclination. IC~342: \\citet[][$I$-band]{wb03}; M74: \\citet[][$R$-band]{m04}; M83: \\citet[][$R$-band]{k+00}; M94: \\citet[][$R$-band]{g+09}; M101: \\citet[][$R$-band extrapolation]{k+00}} \\tablenotetext{f}{From \\citet{sfd98}} \\tablenotetext{g}{IC 342: \\citet{cth00}; M74: \\citet{fb+07}; M83: \\citet{p+01}; M94: \\citet{dB+08}; M101: \\citet{ZEH90}} \\end{deluxetable*}  Planetary nebulae have only recently been used to probe the kinematics of spiral galaxies.  The largest study of this kind has been for the Local Group galaxy M31, where the analysis of over 2000 PN velocities has demonstrated a flattening of the disk's line-of-sight velocity dispersion at large galactocentric radii \\citep{M31PNS}.  Although the high-inclination of the galaxy precluded any dynamical measurement of disk mass, these data did present evidence for disk flaring at $R \\gtrsim 11$~kpc.  The analysis of 140~PN velocities in the moderately inclined ($i \\sim 56^\\circ$) Local Group spiral M33 yielded a similarly surprising result:  by assuming a constant scale-height for the galaxy's disk, \\citet{M33PNe} concluded that the disk mass-to-light ratio actually increases with radius, going from $\\Upsilon_V \\sim 0.3$ in the inner regions to $\\Upsilon_V \\sim 2.0$ at $\\sim 9$~kpc.  Finally, \\citet{M94PNS} measured the velocities of 67 PNe in the low-inclination ($i \\sim 35^\\circ$) spiral M94, and found a line-of-sight velocity dispersion that declines exponentially over $\\sim$4~disk scale lengths.  Papers~I and II of this series \\citep{thesis1, thesis2} presented narrow-band imaging and follow-up spectroscopy of PNe in a sample of large ($r > 7\\arcmin$), nearby ($D < 10$~Mpc), low-inclination ($i \\leq 41^\\circ$) spiral galaxies.  Here we use these data to understand the disk kinematics of five galaxies:  IC~342, M74, M83, M94, and M101.  We introduce our sample in \\S2, and in \\S3, we discuss the problem of the disk scale height, the one parameter of our analysis that cannot be measured.   We explore its dependence on galactic radius, its relationship to other disk properties, such as scale length, and alternatives to the isothermal disk assumption.  We then begin our analysis in \\S4 by compensating for the fact that none of our nearby galaxies is precisely face-on.  We use H{\\sc i} velocity maps to deproject the galactic disks, define the circular velocity as a function of radius, and estimate asymmetric drift.  Next, in \\S5 we identify those PNe which likely belong to the galaxies' spheroidal components, and eliminate them from the analysis.  We find that very few objects fall into this category: only three in IC~342, four in M74, six in M83, three in M94, and one in M101.  Then, in \\S6, we examine the PNe's line-of-sight dispersions without these halo objects.   We show that in the inner regions of our galaxies, the line-of-sight velocity dispersion, $\\sigma_{los}$, generally declines as expected for a constant $h_z$, constant $\\Upsilon$ disk.  However, once past $\\sim$4~optical scale lengths, $\\sigma_{los}$ flattens out at a value much larger than expected.   We consider a few possible explanations for this phenomenon, including the presence of emission-line contaminants in our sample, the effects of internal extinction, and the behavior of large-scale galactic warps.  In \\S7, we model the disks' velocity ellipsoids, and extract $\\sigma_z$ from the line-of-sight velocity dispersions.  As expected, we find that $\\sigma_R$ and $\\sigma_{\\phi}$ contribute little to the measured dispersion of our nearly face-on disks.  In \\S8, we describe the results for each of the five galaxies in our study.  We convert our values of $\\sigma_z$ into disk surface masses and calculate disk mass-to-light ratios.  Then, in \\S9, we analyze the results from the inner and outer regions of the galaxies.  We compute disk and halo rotation curves, compare our disk mass surface densities to those expected from maximal disks, show that pseudo-isothermal cores fit the residual rotation curves better than NFW models, and discuss the physical mechanisms that may produce a constant $\\sigma_z$, including the interaction of thin and thick disks with halo substructure.   We conclude in \\S10 by summarizing our results and suggesting future observations that will assist in their interpretation. ",
        "conclusions": "An analysis of 550 precise PN velocities has yielded kinematic mass estimates throughout the disks of five low-inclination spiral galaxies.  In the inner regions of IC~342, M74, and M94, we find that the velocity dispersion perpendicular to the galactic disk exponentially decays with a scale length twice that of the optical light.  While our observations of the remaining two galaxies yield ambiguous results (due to a Lindbland resonance in M83 and a poorly sampled velocity field in M101), the data are consistent with expectations from a constant scale height, constant mass-to-light ratio disk.   In the later-type (Scd) systems, these disks are clearly sub-maximal, with surface mass densities less than a quarter that needed to reproduce the central rotation curves.  In earlier (Sc) galaxies, this discrepancy is smaller, but still present; only the very-early Sab system M94 has evidence for a maximal disk.  When we subtract off the disk contribution from the H{\\sc i} rotation curve, we infer the potential of the galaxies' invisible dark halos in all our galaxies except M94.  (The uncertainties in the contribution of M94's bright bulge preclude such an analysis in this earliest galaxy in our sample.)   In all four cases, these halos are best fit with a simple pseudo-isothermal model; NFW profiles are poor fits to the data, as they always overpredict the residual rotation curve.  As in the analysis of low surface-brightness and dwarf galaxies, we find no evidence for central dark matter cusps in our systems.  Interestingly, once outside of $\\sim$3~disk scale lengths, the $z$-direction velocity dispersion of galactic disks no longer declines, but instead stays constant with radius.  This behavior, coupled with stability arguments, suggests that the disks of spiral galaxies flare substantially at extremely large radii, and may have strongly increasing mass-to-light ratios.   Whether this flaring is related to the increased importance of the thick disk, or heating of the thin disk by halo substructure, is unclear at the present time.  To further investigate this phenomenon, deep surface photometry of the extreme outer disks of edge-on galaxies is needed.  Our kinematic mass estimates for galactic disks are consistent with expectations.  The galaxies with the latest Hubble types (IC~342, M101) have the smallest mass-to-light ratios, while earlier systems (M74 and M94) have values of $\\Upsilon$ greater than one.  Unfortunately, at this time, no deep, multi-band optical and IR surface photometry exists for any of our galaxies.  Once these data are acquired, we will be able to compare our kinematic values for $\\Upsilon$ to stellar mass-to-light ratio estimates from population synthesis models.  Such a confrontation will not only produce a better constraint on galactic populations, but also will improve our knowledge of the phase-space distribution of galactic disks."
    },
    "0910/0910.0309.txt": {
        "abstract": "New $^{12}$CO $J$=4--3 and $^{13}$CO $J$=3--2 observations of the N159 region, an active site of massive star formation in the Large Magellanic Cloud, have been made with the NANTEN2 and ASTE sub-mm telescopes, respectively. The $^{12}$CO $J$=4--3  distribution is separated into three clumps, each associated with N159W, N159E and N159S. These new measurements toward the three clumps are used in coupled calculations of molecular rotational excitation and line radiation transfer, along with other transitions of the $^{12}$CO $J$=1--0, $J$=2--1, $J$=3--2, and $J$=7--6 as well as the isotope transitions of $^{13}$CO $J$=1--0, $J$=2--1, $J$=3--2, and $J$=4--3. The $^{13}$CO $J$=3--2 data are newly taken for the present work. The temperatures and densities are determined to be $\\sim$ 70-80K and $\\sim 3 \\times 10^3$ cm$^{-3}$ in N159W and N159E and $\\sim$ 30K and $\\sim 1.6 \\times 10^3$ cm$^{-3}$ in N159S. Observed $^{12}$CO $J$=2--1 and $^{12}$CO $J$=1--0 intensities toward N159W and N159E are weaker than expected from calculations of uniform temperature and density, suggesting that low-excitation foreground gas causes self-absorption. These results are compared with the star formation activity based on the data of young stellar clusters and H\\emissiontype{II} regions as well as the mid-infrared emission obtained with the Spitzer MIPS. The N159E clump is associated with embedded cluster(s) as observed at 24 $\\micron$ by the Spitzer MIPS and the derived high temperature, 80K, is explained as due to the heating by these sources.  The N159E clump is likely responsible for a dark lane in a large H\\emissiontype{II} region by the dust extinction. On the other hand, the N159W clump is associated with embedded clusters mainly toward the eastern edge of the clump only. These clusters show offsets of 20$\\arcsec$ - 40$\\arcsec$ from the $^{12}$CO $J$=4--3 peak and are probably responsible for heating indicated by the derived high temperature, 70 K. The N159W clump exhibits no sign of star formation toward the $^{12}$CO $J$=4--3 peak position and its western region that shows enhanced $R_{4-3/1-0}$ and $R_{3-2/1-0}$ ratios. We therefore suggest that the N159W peak represents a pre-star-cluster core of $\\sim 10^5 M_{\\odot}$ which deserves further detailed studies. We note that recent star formation took place between N159W and N159E as indicated by several star clusters and H\\emissiontype{II} regions, while the natal molecular gas toward the stars have already been dissipated by the ionization and stellar winds of the OB stars. The N159S clump shows little sign of star formation as is consistent with the lower temperature, 30K, and somewhat lower density than N159W and N159E. The N159S clump is also a candidate for future star formation. ",
        "introduction": "Giant molecular clouds (hereafter GMCs) are the principal sites of star formation and studies of GMCs are important in understanding the evolution of galaxies. A few tens of GMCs in the solar neighborhood such as Orion A have been well studied, while most of the GMCs in the Galaxy are located at more than a few kpc away in the Galactic disk, where contamination by un-related components in the same line of sight seriously limits detailed understanding of GMCs and their associated objects.  Recent CO surveys of molecular clouds toward external galaxies in the Local Group have revealed that properties of GMCs such as relations between CO luminosity and line width, the mass range, and the index of mass spectrum are similar among these galaxies \\citep{2007prpl.conf...81B}. This suggests that the properties of GMCs are fairly common among galaxies. The Large Magellanic Cloud (hereafter LMC) is the nearest galaxy to our own at a distance of 50 kpc and is nearly face on, making it an ideal laboratory to observe various properties of GMCs. The small distance enables us to make molecular observations at high spatial resolutions. The LMC also provides a unique opportunity to study molecular clouds and star formation in different environments from the Galaxy; the gas-to-dust ratio is $\\sim 4$ times higher \\citep{1982A&A...107..247K}, and the metal abundance is about $\\sim $3 -- 4 times lower \\citep{2002A&A...396...53R,1984IAUS..108..353D} than those of the Galaxy.  Following a low resolution $^{12}$CO $J$=1--0 survey at 140 pc resolution with the Columbia 1.2 m telescope located at CTIO in Chile \\citep{1988ApJ...331L..95C}, high resolution observations in the $^{12}$CO $J$=1--0, $J$=2--1 emission line were made with the SEST 15 m telescope toward some of the molecular clouds and revealed their clumpy structure at 10 pc resolution (e.g., \\cite{1986ApJ...303..186I,1994A&A...291...89J,1996ApJ...472..611C,1997A&AS..122..255K,1998A&A...331..857J,2003A&A...406..817I}). However, these high resolution studies were limited in the spatial coverage, compared with the large angular extent of the LMC, $\\sim 6\\arcdeg \\times 6\\arcdeg$. \\citet{2008ApJS..178...56F} carried out a survey in the $^{12}$CO $J$=1--0 emission line at 40 pc resolution over a $6\\arcdeg \\times 6\\arcdeg$ field in the LMC with the NANTEN 4 m telescope and obtained a complete sample of 270 GMCs (see also for preceding works \\cite{1999PASJ...51..745F,2001PASJ...53L..41F,2001PASJ...53..971M,2001PASJ...53..985Y}). These studies revealed that young stellar clusters whose ages are less than 10 Myr are spatially well correlated with GMCs, and that GMCs are categorized into three types in terms of star formation activity. Type I is starless in the sense that they are not associated with O stars, Type II is associated with small H\\emissiontype{II} regions only, and Type III is associated with huge H\\emissiontype{II} regions and young stellar clusters \\citep{Kawamura09, 2007prpl.conf...81B}. These types are interpreted in terms of evolutionary sequence of GMCs in a timescale of 2-30 Myrs \\citep{1999PASJ...51..745F,2007prpl.conf...81B,Kawamura09}.  The $^{12}$CO $J$=1--0 emission line is a probe commonly used to trace molecular clouds because of its low excitation energy ($\\sim$ 5K) and low critical density for collisional excitation ($n_{cr} \\sim 1000$cm$^{-3}$). The $^{12}$CO $J$=1--0 emission alone is however not able to provide physical properties like kinetic temperature and density, the fundamental parameters of GMCs. The high $J$ transitions have higher excitation energies and higher critical densities; e.g., the $^{12}$CO $J$=3--2 transition has the upper state at 33 K and the critical density, $3\\times 10^4 $cm$^{-3}$ \\citep{2005A&A...432..369S}, and the $^{12}$CO $J$=4--3 transition has the upper state at 55 K and the critical density, $1\\times 10^5 $cm$^{-3}$. These sub--mm CO emission lines can selectively trace the sites which may be warmer and denser than the mm CO lines and have a potential to reveal detailed physical properties where star formation is taking place. We are also allowed to make better estimates of temperatures and densities of molecular clouds with a combination of multi--$J$ CO line intensities and molecular excitation analyses such as the Large Velocity Gradient (hereafter LVG) model of molecular line transfer. Recently-developed sub--mm telescopes such as ASTE, NANTEN2 and APEX located at altitudes around 5000m in Atacama, Chile have enabled us to observe higher excited CO transitions at sub--mm wavelengths in superb observational conditions \\citep{2004SPIE.5489..763E,2006IAUSS...1E..21F,2006A&A...454L.115G}. At a lower angular resolution, AST/RO also offered sub-mm observing capability \\citep{2001PASP..113..567S}.  \\citet{2000ApJ...545..234B} first detected the $^{12}$CO $J$=4--3 emission line toward the N159 region at 50 pc resolution with the AST/RO telescope. \\citet{2005ApJ...633..210B} later presented estimates of temperatures and densities; $T_{\\rm kin}$ = 20 K, $n({\\rm H}_2)$ = 10$^5$ cm$^{-3}$ for the cold dense component and $T_{\\rm kin}$ = 100 K, $n({\\rm H}_2)$ = 100 cm$^{-3}$ for hot tenuous component toward N159W. Subsequently, \\citet{2008ApJS..175..485M} carried out high resolution $^{12}$CO $J$=3--2 observations of several GMCs at 5 pc resolution with the ASTE 10 m telescope including N159 in the LMC. \\citet{2008ApJS..175..485M} revealed detailed structure of highly excited gas with $T_{\\rm kin} \\sim$ 20 -- 200 K and $n({\\rm H}_2)$ $\\sim$ 10$^3$ -- 10$^4$ cm$^{-3}$. In the present study, we shall focus on the N159 GMC which shows the highest $^{12}$CO $J$=1--0 intensity from the NANTEN survey \\citep{2008ApJS..178...56F}. The N159 GMC is classified as Type III, and includes at least three prominent clumps, N159W, N159E and N159S \\citep{1994A&A...291...89J}. The two molecular clumps in the northern part, N159W and N159E, are associated with massive star clusters whose ages are younger than 10 Myr \\citep{1996ApJS..102...57B} and with huge H\\emissiontype{II} regions \\citep{1956ApJS....2..315H,1976MmRAS..81...89D,1986ApJ...306..130K}. On the other hand, there is no star formation in N159S. In order to better estimate the physical properties of the N159 region, we have carried out new sub-mm observations in the $^{12}$CO $J$=4--3 emission line at a 10 pc spatial resolution with NANTEN2 and in the $^{13}$CO $J$=3--2 emission line at a 5 pc spatial resolution with ASTE. We have also made combined calculations of molecular rotational excitation and line radiative transfer by employing the CO datasets available in N159 and compared the results with the star formation activity.  The present paper is organized as follows: Section \\ref{pasj_obs} describes the observations. Sections \\ref{pasj_res} and \\ref{pasj_analysis} show the observational results and data analysis, respectively. We discuss the correlation between highly excited molecular gas and star formation activities in section \\ref{pasj_sf}. Finally, we present a summary in section \\ref{pasj_sum}. ",
        "conclusions": "\\label{pasj_sum} We carried out $^{12}$CO $J$=4--3 observations of the N159 region with NANTEN2 and $^{13}$CO $J$=3--2 observations toward $^{12}$CO $J$=3--2 peaks with ASTE. The main results are summarized as follows;  \\begin{enumerate} \\item The N159 GMC has been resolved into three prominent clumps, N159W, N159E, and N159S in the $^{12}$CO $J$=4--3 emission line. N159W shows the highest intensities among the three in the CO transitions from $J$=1--0 to $J$=7--6.  \\item Molecular densities and temperatures have been derived toward the three peaks. Using a LVG analysis involving the $^{12}$CO $J$=4--3 , $^{13}$CO $J$=3--2 emission lines in addition to the other CO lines published, we find that N159W and N159E have temperature of $\\sim$ 70 -- 80 K and density of $\\sim 3 \\times 10^3$ cm$^{-3}$, and that N159S has temperature of $\\sim$ 30 K and density of $\\sim 1.6 \\times 10^3$ cm$^{-3}$. In order to explain lower line intensities than expected, we suggest that $^{12}$CO $J$=1--0 and $J$=2--1 lines may be affected by self-absorption by foreground lower excitation gas.  \\item The $^{12}$CO $J$=4--3 distribution is compared with H$\\alpha$ and mid-- to far--infrared emission, a sign of embedded star formation, obtained with the Spitzer SAGE program. N159E shows a good coincidence with a dark lane of H$\\alpha$ and also with a 24 $\\micron$ extended source. On the other hand, N159W is associated with three compact 24 $\\micron$ sources and some small H$\\alpha$ features, although the $^{12}$CO $J$=4--3 peak of N159W and its western part show no sign of star formation.  \\item A comparison between N159 and $\\eta$ Car indicates that they show similar star formation activity and we do not see significant difference in physical parameters between these two massive star forming regions. \\end{enumerate}  \\bigskip We thank the all members of the NANTEN2 consortium and ASTE team for the operation and persistent efforts to improve the telescopes. This research was supported by the Grant-in-Aid for Nagoya University Global COE Program, \"Quest for Fundamental Principles in the Universe: from Particles to the Solar System and the Cosmos\", from the Ministry of Education, Culture, Sports, Science and Technology of Japan. This work is financially supported in part by a Grant-in-Aid for Scientific Research from the Ministry of Education, Culture, Sports, Science and Technology of Japan (No. 15071203) and from JSPS (Nos. 14102003 and 18684003), and by the JSPS core-to-core program (No. 17004), and the Mitsubishi Foundation. This work is also financially supported in part by the grant SFB 494 of the Deutsche Forschungsgemeinschaft, the Ministerium fur Innovation, Wissenschaft, Forschung und Technologie des Landes Nordrhein-Westfalen and through special grants of the Universitat zu Koln and Universitat Bonn. SAGE research has been funded by NASA/Spitzer grant 1275598 and NASA NAG5-12595. The ASTE project is driven by Nobeyama Radio Observatory (NRO), a branch of National Astronomical Observatory of Japan (NAOJ), in collaboration with University of Chile, and Japanese institutes including University of Tokyo, Nagoya University, Osaka Prefecture University, Ibaraki University., and Hokkaido University. Observations with ASTE were in part carried out remotely from Japan by using NTT's GEMnet2 and its partner R\\&E (Research and Education) networks, which are based on AccessNova collaboration of University of Chile, NTT Laboratories, and NAOJ. A part of this study was financially supported by the MEXT Grant-in-Aid for Scientific Research on Priority Areas No.\\ 15071202."
    },
    "0910/0910.5219_arXiv.txt": {
        "abstract": "Measurements of X-ray scaling laws are critical for improving cosmological constraints derived with the halo mass function and for understanding the physical processes that govern the heating and cooling of the intracluster medium. In this paper, we use a sample of \\ngroup X-ray selected galaxy groups to investigate the scaling relation between X-ray luminosity ($\\rm{L}_{\\rm X}$) and halo mass ($\\rm{M}_{\\rm 200}$) where $\\rm{M}_{\\rm 200}$ is derived via stacked weak gravitational lensing. This work draws upon a broad array of multi-wavelength COSMOS observations including 1.64 degrees$^2$ of contiguous imaging with the Advanced Camera for Surveys (ACS) to a limiting magnitude of $\\rm{I}_{ \\rm F814W}=26.5$ and deep {\\sl XMM-Newton/Chandra} imaging to a limiting flux of $1.0\\times 10^{-15}$ \\flux in the 0.5-2 keV band. The combined depth of these two data-sets allows us to probe the lensing signals of X-ray detected structures at both higher redshifts and lower masses than previously explored. Weak lensing profiles and halo masses are derived for nine sub-samples, narrowly binned in luminosity and redshift. The COSMOS data alone are well fit by a power law, ${\\rm M}_{\\rm 200} \\propto (\\rm{L}_{\\rm X})^{\\alpha}$, with a slope of \\cosalpha . These results significantly extend the dynamic range for which the halo masses of X-ray selected structures have been measured with weak gravitational lensing. As a result, tight constraints are obtained for the slope of the \\mlx relation. The combination of our group data with previously published cluster data demonstrates that the \\mlx relation is well described by a single power law, \\henkalpha, over two decades in mass, $\\rm{M}_{\\rm 200} \\sim 10^{13.5}$ -- $10^{15.5} \\mass$. These results are inconsistent at the \\sssig\\ level with the self-similar prediction of $\\alpha=0.75$. We examine the redshift dependence of the \\mlx relation and find little evidence for evolution beyond the rate predicted by self-similarity from $z \\sim 0.25$ to $z \\sim 0.8$. ",
        "introduction": "Groups and clusters of galaxies, formed through the gravitational collapse of massive dark matter halos, are now readily identified up to redshift one and even beyond \\citep[\\eg][]{Stanford:2006,Eisenhardt:2008}. Baryonic tracers such as red-sequence galaxies, typically abundant at the centers of groups and clusters, or X-ray emission from the hot intracluster medium (ICM), have proved to be especially successful in this task \\citep[\\eg][]{Gladders:2005,Koester:2007,Finoguenov:2007}. Nonetheless, these observables only trace the tip of the iceberg given that the vast majority of the underlying mass is in the form of dark matter.  The quantification of the total mass (both dark and baryonic) of groups and clusters of galaxies is an important endeavour from both a cosmological and an astronomical standpoint. In particular, several lines of research would benefit from a clearer understanding of the relationship between the total halo mass of groups and clusters and their baryonic tracers. We outline several briefly here \\citep[for a recent review on this subject see][]{Voit:2005}. From a cosmological perspective, the number density of groups and clusters as a function of total mass is of fundamental interest because it is sensitive to both the expansion and growth history of the universe and can be used to constrain cosmological parameters such as $\\Omega_m, \\sigma_8,$ and $\\Omega_{\\Lambda}$ \\citep[\\eg][]{White:1993,Wang:1998a,Haiman:2001,Rozo:2004, Wang:2004a, Bahcall:2004, Rozo:2009}. Furthermore, modifications to the laws of gravity which can be invoked as a possible physical mechanism for acceleration, could imprint telltale signatures in the abundance and dark matter structure of groups and clusters of galaxies \\citep[][]{Rapetti:2008,Schmidt:2008}. Unfortunately, the common difficulty encountered with each of these enquires is that theories and simulations make dark matter based predictions but our most accessible observables (such as richness or X-ray luminosity) are baryonic in nature.  It has long been recognized that baryonic observables are subject to complex and poorly understood physical processes that make them imperfect dark matter tracers. For example, X-ray studies discovered early on that the theory of pure gravitational collapse which makes simple predictions for the shape and amplitude of X-ray scaling relations \\citep[also known as the {\\it self-similar model,}][]{Kaiser:1986}, fails to match observations such as the slope and normalization of the \\lx-T relation \\citep[][and references therein]{Voit:2005} implying that other non-gravitational (and still much debated) processes have significantly affected the thermodynamic state of the ICM. Additional heating and cooling mechanisms that are invoked to solve this puzzle lead to different predictions for the shape and redshift evolution of X-ray scaling relations. On the one hand, the fact that X-ray scaling relations deviate from simple models is a plague for cosmologists because there is no straightforward recipe to estimate total halo masses. On the other hand, from an astronomical perspective, the comparison between X-ray observables and total halo mass contains valuable clues about the physical processes that govern galaxy formation and the heating and cooling of the ICM.  For all of these reasons, more precise measurements of the mean and scatter in the relationship between total halo mass and various baryonic tracers of groups and clusters of galaxies are highly desirable (\\eg\\ see discussions in \\citet{Voit:2005} and \\citet{Albrecht:2006}).  At present, there are five popular methods for detecting groups and clusters of galaxies: a) optical detection via the red-sequence \\citep[\\eg][]{Gladders:2005,Koester:2007}, b) detection via the Sunyaev-Zeldovich (SZ) effect which measures the distortion of the CMB spectrum due to the hot ICM \\citep[\\eg][]{Sunyaev:1970,Sunyaev:1972,Carlstrom:2002,Benson:2004,Staniszewski:2008}, c) detection via X-ray emission \\citep[\\eg][]{Bohringer:2000,Hasinger:2001,Finoguenov:2007, Vikhlinin:2008}, d) spectroscopic identification \\citep[\\eg][]{Gerke:2005,Miller:2005,Knobel:2009}, and e) detection via weak lensing maps \\citep[\\eg][]{Marian:2006,Miyazaki:2007,Massey:2007a}. This last technique is the simplest in terms of the underlying physics and is the only method for which the total halo mass can be directly probed, independently of both the baryons and the dynamical state of the cluster. However, shear maps can only detect the most massive systems ($M>10^{14} M_{\\odot}$) and are limited to moderate redshifts because the lensing weight function peaks mid-way between the source and the observer, with galaxy shapes increasingly difficult to measure at $z>1$. X-ray observations on the other hand, can more simply probe complete samples of groups and clusters, but departures from virial equilibrium and non-thermal pressure components in the ICM can bias X-ray based hydrostatic mass estimates \\citep[\\eg][]{Nagai:2007}. The SZ effect has the attractive property of being redshift independent, and the integrated SZ flux increment, Y, may be less sensitive to the baryon physics of cluster cores \\citep[][]{Motl:2005,Nagai:2006} but mass measurements with the SZ effect face other challenges such as the identification and removal of radio point-sources \\citep[][]{Vale:2006}, sky confusion owing to projection effects \\citep{White:2002a}, and possibly a larger scatter in the Y-M relation than previously estimated due to feedback processes \\citep[][]{Shaw:2008}.  Given these considerations, a promising strategy is to employ a robust and efficient cluster detection method \\citep[to which the ultimate solution may be a combination of several techniques such as described in][]{Cohn:2009} and to perform an absolute mass calibration of baryonic tracers via weak gravitational lensing \\citep[][]{Hoekstra:2007a,Rykoff:2008,Berge:2008}.  The focus of this paper is to advance these goals by calibrating the slope and amplitude of the \\mlx relation for galaxy groups using cross-correlation weak lensing in the COSMOS survey (also called ``group-galaxy lensing''). Extending weak lensing measurements into the group regime is particularly important in order to extend the dynamic range of weak-lensing-based mass-estimates so as to more accurately determine the slopes of scaling relations. Using the COSMOS sample, we show that X-ray detections span a more complete and wide range of redshift and mass than detections via shear maps. Indeed, high redshift and small structures are challenging to detect directly with shear because of the shape of the lensing weight function (see $\\S$\\ref{group_selection}). Nevertheless, once they have been identified, groups and clusters can be studied via stacking techniques. A notable advantage of this method is that measurements are unaffected by uncorrelated mass along the line-of-sight whereas mass estimates for individually detected clusters are subject to $\\sim 20 \\%$ uncertainties \\citep[][]{Metzler:2001,Hoekstra:2003a,de-Putter:2005}. The associated drawback with stacking is that the intrinsic scatter around the mean relation is difficult to recover.  In order to employ the stacked weak lensing technique, tight baryonic tracers of halo mass are highly desirable. The X-ray luminosity of groups and clusters is considered to be a reasonable tracer of halo mass with a logarithmic scatter in the \\mlx relation of roughly $20\\%$ to $30\\%$ \\citep[][]{Stanek:2006, Maughan:2007, Pratt:2008, Rozo:2008, Rykoff:2008, Vikhlinin:2009}. A large fraction of this scatter has been shown to be associated with the presence of cool-cores and simple excision techniques can reduce the scatter to sub-$20\\%$ levels \\citep[][]{Maughan:2007,Pratt:2008}. Although more tightly correlated mass tracers have been identified such as the ${\\rm Y_{X}}$ parameter \\citep[\\eg][]{Kravtsov:2006} -- such indicators require the measurement of an X-ray spectrum which is not possible for most survey data where count rates are low. Our choice of \\lx as a mass proxy reflects that fact that it is a simple X-ray observable, accessible with survey quality data, and the only one that can be easily measured at high redshift. Temperature measurements may be feasible for a small fraction of high redshift objects but cosmological studies that require complete samples of high redshift groups and clusters will need simple mass proxies like ${\\rm{L_{\\rm X}}}$. The details of the \\mlx relation are also important (regardless of the choice of a mass proxy) for determining effective volumes as a function of mass in X-ray flux limited surveys \\citep[][]{Stanek:2006, Vikhlinin:2009}.  The layout of this paper is as follows. The data are presented in $\\S$\\ref{cosmos_survey} and the theoretical lensing background is developed in $\\S$\\ref{lensing_theory}. The construction of the group catalog and the lens selection are specified in $\\S$\\ref{group_selection}. Details regarding the adopted form of the \\mlx relation are given in $\\S$\\ref{lx_m_relation}. Our main results are presented in $\\S$\\ref{results} followed by our assessment of the systematic errors in $\\S$\\ref{syst_error}. Finally, we discuss the results and draw up our conclusions in $\\S$\\ref{discussion}.  We assume a WMAP5 $\\Lambda$CDM cosmology with $\\Omega_{\\rm m}=0.258$, $\\Omega_\\Lambda=0.742$, $\\Omega_{\\rm b}h^2=0.02273$, $n_{\\rm s}=0.963$, $\\sigma_{8}=0.796$, $H_0=72$ $h_{72}$ km~s$^{-1}$~Mpc$^{-1}$ \\citep[][]{Hinshaw:2009}. All distances are expressed in physical units of $h_{72}^{-1}$ Mpc. X-ray luminosities are expressed in the 0.1-2.4 keV band, rest-frame. The letter M denotes halo mass in general whereas \\m is explicitly defined as $M_{200}\\equiv M(<r_{200})=200\\rho_{crit}(z) \\frac{4}{3}\\pi r_{200}^3$ where $r_{200}$ is the radius at which the mean interior density is equal to 200 times the critical density ($\\rho_{crit}\\equiv 3H^2(z)/8\\pi G$). The function $E(z) \\equiv H(z)/H_0 = \\sqrt{\\Omega_{m}(1+z)^3+\\Omega_{\\Lambda}}$ represents the Hubble parameter evolution for a flat metric. All magnitudes are given on the AB system. ",
        "conclusions": ""
    },
    "0910/0910.4735_arXiv.txt": {
        "abstract": "Although originally conceived as primarily an extragalactic survey, the Sloan Digital Sky Survey (SDSS-I), and its extensions SDSS-II and SDSS-III, continue to have a major impact on our understanding of the formation and evolution of our host galaxy, the Milky Way. The sub-survey SEGUE: Sloan Extension for Galactic Exploration and Understanding, executed as part of SDSS-II, obtained some 3500 square degrees of additional $ugriz$ imaging, mostly at lower Galactic latitudes, in order to better sample the disk systems of the Galaxy. Most importantly, it obtained over 240,000 medium-resolution spectra for stars selected to sample Galactocentric distances from 0.5 to 100 kpc. In combination with stellar targets from SDSS-I, and the recently completed SEGUE-2 program, executed as part of SDSS-III, the total sample of SDSS spectroscopy for Galactic stars comprises some 500,000 objects.  The development of the SEGUE Stellar Parameter Pipeline has enabled the determination of accurate atmospheric parameter estimates for a large fraction of these stars. Many of the stars in this data set within 5 kpc of the Sun have sufficiently well-measured proper motions to determine their full space motions, permitting examination of the nature of much more distant populations represented by members that are presently passing through the solar neighborhood. Ongoing analyses of these data are being used to draw a much clearer picture of the nature of our galaxy, and to supply targets for detailed high-resolution spectroscopic follow-up with the world's largest telescopes. Here we discuss a few highlights of recently completed and ongoing investigations with these data. ",
        "introduction": "The era of the massive surveys of Galactic stars is now very much underway. Our understanding of the history of the formation and evolution of the stellar populations in the Galaxy is presently being revolutionized, based on results from two primary surveys, the Sloan Digital Sky Survey (SDSS; \\cite[York et al. 2000]{York00}), in particular its dedicated sub-surveys Sloan Extension for Galactic Exploration and Understanding (SEGUE-1 and SEGUE-2; \\cite[Yanny et al. 2009]{Yan09}) and the RAdial Velocity Experiment (RAVE; \\cite[Zwitter et al. 2008]{Zwi08}; Wyse, this volume). Collectively, these two surveys are in the process of providing detailed spectroscopic information for a million or more stars, sampling all of the known stellar populations in the Galaxy. Although refinements in the procedures for extracting estimates of the stellar atmospheric parameters (T$_{\\rm eff}$, log g, [Fe/H]) and individual element abundances (in the case of SDSS/SEGUE, [$\\alpha$/Fe] and [C/Fe] ratios; for RAVE, many more) are still underway, the wealth of information already available provides the basis for numerous investigations and follow-up observations.\\footnote{The most recent public release for SDSS/SEGUE is DR-7, which includes SEGUE-1; see \\cite[Abazajian et al. (2009)]{Aba09}. The next public release for SDSS/SEGUE is DR-8, scheduled for December 2010, which will include results from SEGUE-2.}  Here we summarize a few of the ongoing projects that are being enabled by the SDSS/SEGUE database.  \\begin{figure}[t] \\begin{center} \\includegraphics[width=4.0in]{beers_fig1.eps} \\caption{Simple visualization of the SDSS/SEGUE stellar database, for roughly 400,000 stars with available atmospheric parameter estimates from the SSPP.  The sphere has a radius of 10 kpc, centered on the Sun.  The colors scale with metallicity, from green (metal-poor stars) to red (metal-rich stars).} \\label{fig1} \\end{center} \\end{figure} ",
        "conclusions": ""
    },
    "0910/0910.1110_arXiv.txt": {
        "abstract": "We report results from a 2007 {\\it Suzaku} observation of the Seyfert 1 AGN NGC 4593. The narrow Fe K$\\alpha$ emission line has a FWHM width $\\sim$ 4000 km s$^{-1}$, indicating emission from $\\gtrsim$ 5000 $R_{\\rm g}$. There is no evidence for a relativistically broadened Fe K line, consistent with the presence of a radiatively efficient outer disk which is truncated or transitions to an interior radiatively inefficient flow.  The {\\it Suzaku} observation caught the source in a low-flux state; comparison to a 2002 {\\it XMM-Newton} observation indicates that the hard X-ray flux decreased by 3.6, while the Fe K$\\alpha$ line intensity and width $\\sigma$ each roughly halved. Two model-dependent explanations for the changes in Fe K$\\alpha$ line profile are explored. In one, the Fe K$\\alpha$ line width has decreased from $\\sim$10000 to $\\sim$4000 km s$^{-1}$ from 2002 to 2007, suggesting that the thin disk truncation/transition radius has increased from 1000--2000 to $\\gtrsim$5000 $R_{\\rm g}$. However, there are indications from other compact accreting systems that such truncation radii tend to be associated only with accretion rates relative to Eddington much lower than that of NGC 4593. In the second model, the line profile in the {\\it XMM-Newton} observation consists of a time-invariant narrow component plus a broad component originating from the inner part of the truncated disk ($\\sim$300 $R_{\\rm g}$) which has responded to the drop in continuum flux. The Compton reflection component strength $R$ is $\\sim$ 1.1, consistent with the measured Fe K$\\alpha$ line total equivalent width with an Fe abundance 1.7 times the solar value. The modest soft excess, modeled well by either thermal bremsstrahlung emission or by Comptonization of soft seed photons in an optical thin plasma, has fallen by a factor of $\\sim$20 from 2002 to 2007, ruling out emission from a region 5 lt-yr in size. ",
        "introduction": "Accretion onto supermassive black holes in many Active Galactic Nuclei (AGN) is generally thought to proceed via a radiatively efficient, optically thick, geometrically thin disk (``$\\alpha$-disk'', e.g., Shakura \\& Sunyaev 1973), as evidenced by their optical/UV continua (the ``big blue bump''; Sun \\& Malkan 1989). Galactic black hole systems also contain evidence for such a component; the thermal blackbody emission extends into the soft X-ray band. In Seyferts, the Fe K$\\alpha$ emission line at 6.4 keV line is a key tracer of the radiatively efficient circumnuclear material. Narrow lines with Doppler-broadened FWHMs of a few thousand km s$^{-1}$ are virtually ubiquitous in X-ray spectra of low-redshift Seyferts observed with {\\it XMM-Newton}, {\\it Chandra}-HETGS or {\\it Suzaku} (e.g., Yaqoob \\& Padmanabhan 2004, Nandra 2006) and indicate material originating light-days or farther from the black hole, e.g., in the outer disk or molecular torus. Fe K$\\alpha$ line emission originating near the innermost stable orbit of the disk yields a broad (FWHMs $\\sim$ 0.1$c$) and redshifted profile sculpted by general and special relativistic effects in the regime of strong gravity (Fabian et al.\\ 1989, 2002), although to accurately gauge the strength of the line, one must correctly model any continuum curvature which may be associated with ionized absorption along the line of sight (e.g., Reeves et al.\\ 2004, Turner et al.\\ 2005).   However, the flow in compact systems accreting at relatively low values relative to the Eddington limit may contain a radiatively inefficient component (a radiatively inefficient accretion flow or RIAF; Narayan et al.\\ 1998, Quataert 2001); a configuration consisting of an inner RIAF flow which is surrounded by an $\\alpha$-disk beyond some transition radius (e.g., Esin et al.\\ 1997) has been suggested for some low-luminosity AGN (Lu \\& Wang 2000 and references therein). In the case of the Seyfert 1 AGN NGC 4593, modeling of the optical/UV continuum indicates blackbody emission from a truncated thin disk, with an inner radius of emission of 30 $R_{\\rm g}$\\footnote{1  $R_{\\rm g} \\equiv GM_{\\rm BH}/c^2$} (Lu \\& Wang 2000). An observation with {\\it XMM-Newton} in 2002 indicated a FWHM line width of 11000$\\pm$1000 km s$^{-1}$; assuming a black hole mass $M_{\\rm BH}$ of $10^{6.8}$ $\\Msun$, this width suggests an inner extent of no less than roughly 1000 $R_{\\rm g}$ (Brenneman et al.\\ 2007, hereafter B07). No broad Fe K$\\alpha$ line has been confirmed with {\\it XMM-Newton} (Reynolds et al.\\ 2004, B07). These results were consistent with those obtained from a {\\it Chandra}-HETGS observation in 2001, from which Yaqoob \\& Padmanabhan (2004) measured a FWHM line width of $2140^{+13780}_{-1830}$ km s$^{-1}$. An earlier {\\it ASCA} observation in 1994 indicated line emission from a radius of $\\geq$30$R_{\\rm g}$, consistent with the thin disk truncation radius suggested from optical/UV continuum (``big blue bump'') spectral energy distribution (SED) fitting (Lu \\& Wang 2000). Furthermore, the 2002 {\\it XMM-Newton} observation revealed, in addition to the Fe K$\\alpha$ core, line emission from ionized Fe, likely \\ion{Fe}{26} (B07). Such line features are relatively rare in Seyfert X-ray spectra, and could potentially originate in the collisionally ionized transition region between the inner radiatively inefficient flow and the thin disk (e.g., Lu \\& Wang 2000).  Another diagnostic reflection component present in hard X-ray spectra is the Compton reflection hump peaking at $\\sim$20--30 keV, expected when the hard X-ray power-law continuum illuminates optically thick material. The strength of the Compton reflection hump $R$ was found to be $\\sim$ 1 in a 1997 observation with {\\it BeppoSAX} (Guainazzi et al.\\ 1999); $R$ = 1 corresponds to a slab covering 2$\\pi$ sr of the sky as seen from the illuminating X-ray source. The {\\it Suzaku} X-ray observatory is the first since {\\it BeppoSAX} to provide continuous coverage from below 1 keV to at least 50 keV, allowing users to spectrally deconvolve the broadband continuum components (absorbing components, the primary power-law, and the Compton reflection hump), but {\\it Suzaku}'s lower $>$10 keV background yields a more accurate determination of $R$.  In this paper, we report results from an observation of the nucleus of NGC 4593 with {\\it Suzaku} in 2007, with the goals of constraining the Fe K emission profile and accurately determining the strength of the Compton reflection hump in order to constrain the geometry of the circumnuclear accreting gas. As demonstrated below, {\\it Suzaku} caught the source at an atypically low 2--10 keV flux level, a factor of almost 4 lower than during the {\\it XMM-Newton} observation. We observe significant changes in the Fe K$\\alpha$ profile between the 2002 {\\it XMM-Newton} and 2007 {\\it Suzaku} observations which may be related to the decrease in continuum flux. We also report evolution in the strength of the soft excess, and, tentatively, the ionized Fe K emission.  The rest of this paper is organized as follows: Section 2 describes the observations and data reduction. In Section 3, we present fits to the Fe K emission complex observed with {\\it Suzaku} and compare the results to those for the 2002 {\\it XMM-Newton} observation. In Section 4, we present fits to the 0.3--76 keV broadband {\\it Suzaku} time-averaged spectrum, and again compare the results to those obtained from {\\it XMM-Newton} to investigate long-term spectral variability of the broadband emission components. The results are discussed in Section 5. ",
        "conclusions": "\\subsection{Summary of Observational Results}  We have presented results from a {\\it Suzaku} observation of the nucleus of the Seyfert AGN NGC 4593 in 2007 December, and we compare our spectral fits for both the Fe K bandpass and the broadband X-ray spectrum with those obtained from a 2002 {\\it XMM-Newton} EPIC-pn observation. {\\it Suzaku} caught the source at a relatively low X-ray flux level: the 2--10 keV continuum flux during the {\\it Suzaku} observation was a factor of 3.8 lower.  The Fe K$\\alpha$ line intensity has dropped by a factor of 1.7, suggesting that roughly half of the total line flux has responded to the drop in continuum flux. One of our main results is that the Fe K$\\alpha$ line is significantly more narrow in the {\\it Suzaku} observation. Modeling the line as a single Gaussian, we find that the width $\\sigma$ has dropped from $87 \\pm 17$ eV in 2002 to $41^{+12}_{-16}$ eV in 2007. We also modeled the line using a dual-Gaussian model composed of relatively narrow and broad lines. The former dominates the {\\it Suzaku} profile and is assumed to be time-invariant; in the {\\it XMM-Newton} spectrum, both lines are modeled to have roughly equal intensity and the broad component has a width $\\sigma = 177^{+84}_{-52}$ eV. There is highly tentative evidence for the \\ion{Fe}{26} emission line at 6.96 keV to have dropped in intensity from 2002 to 2007, assuming an intrinsically narrow (unresolved) line: in the {\\it Suzaku} observation, the line is not significantly detected ($EW < 17$ eV).   In our broadband fits to the 0.3--76 keV spectrum, the primary power-law component, commonly attributed to inverse Comptonization of soft seed photons in a hot corona (e.g., Haardt et al.\\ 1994), was relatively flat, with $\\Gamma$ near 1.65. The Compton reflection component had a relative strength $R \\sim 1.08$. We also model a modest soft excess using both thermal bremsstrahlung emission and thermal Comptonization of soft seed photons, similar to B07, and we obtain similar results. Importantly, we find the soft excess has dropped in flux by a factor of at least $\\sim$20 between the {\\it XMM-Newton} and {\\it Suzaku} observations. We model one zone of absorption along the line of sight, the previously seen highly-ionized (log$\\xi$ $\\sim$ 2.5)  zone, with a column density similar to that obtained by McKernan et al.\\ (2003) and Steenbrugge et al.\\ (2003). There is no strong evidence for evolution of the warm absorber between the two observations. A relativistically broadened Fe K$\\alpha$ line was not detected in the {\\it Suzaku} spectrum; Reynolds et al.\\ (2004) demonstrated a similar result in the {\\it XMM-Newton} spectrum.      \\subsection{Tracing the Truncated Disk with the Fe K$\\alpha$ Line}    %   We explore two (model-dependent) scenarios to correlate the changes in Fe line intensity and profile with the observed drop in continuum flux. In the model where a single Gaussian was used to describe the Fe K$\\alpha$ profile, the width $\\sigma$ in the 2007 {\\it Suzaku} observation was $41^{+12}_{-16}$ eV, which corresponds to a FWHM velocity $v_{\\rm FWHM}$ of $4420^{+1290}_{-1730}$ km s$^{-1}$. This is roughly commensurate with the optical broad emission line width: Peterson et al.\\ (2004) reported FWHM H$\\alpha$ and H$\\beta$ line widths of $3399\\pm196$ and $3769\\pm862$ km s$^{-1}$, respectively.\\footnote{Peterson et al.\\ (2004) reported that the continuum-line lag results were poor. The H$\\alpha$ lag was reported as $3.2^{+5.6}_{-4.1}$ lt-days but Peterson et al.\\ (2004) recommended caution. The H$\\beta$ lag was reportedly ``completely unreliable.''} Assuming that the line originates in gas that is in virialized orbit around the black hole, we can estimate the distance $r$ from the black hole to the line-emitting gas. Assuming that the velocity dispersion is related to $v_{\\rm FWHM}$ as $<$$v^2$$>$ = $\\frac{3}{4}v^2_{\\rm FWHM}$ (Netzer 1990), we use $G$$M_{\\rm BH}$ = $r$$v^2$. We use a black hole mass $M_{\\rm BH}$ of $6.6 \\times 10^6 \\Msun$, an estimate based on the relation between $M_{\\rm BH}$ and stellar velocity dispersion in Seyferts (Nelson et al.\\ 2004). The best-fit reverberation mapping estimate from Peterson et al.\\ (2004), $5.4 \\times 10^6 \\Msun$, is consistent with this estimate.  We find $r$ = $6.0^{+10.2}_{-2.4} \\times 10^{13}$ m, or $2.3^{+3.9}_{-1.0}$ lt-days. As 1 $R_{\\rm g}$ corresponds to $1 \\times 10^{10}$ m for the black hole mass used, $r$ = $6000^{+10500}_{-2500} R_{\\rm g}$. We cannot of course rule out contribution from an even more narrow Gaussian component originating in even more distant material. In the 2002 {\\it XMM-Newton} observation, the corresponding measured line width (we use our best-fit value of $\\sigma = 87 \\pm 17$ eV) corresponds to a value of $r$ = $1.3^{+0.7}_{-0.4} \\times 10^{13}$ m = $0.50^{+0.30}_{-0.15}$ lt-days, or $1350^{+750}_{-350} R_{\\rm g}$ (see also B07). B07 also used the lack of observed variability in the Fe line flux during the {\\it XMM-Newton} observation to constrain the light-crossing time for the line-emitting gas to be at least 2000 $R_{\\rm g}$.  One possible explanation to explain the change in Fe K$\\alpha$ line profile, insofar as it traces the geometrically thin, radiatively efficient disk, is that the innermost radius of the thin disk has increased over 5 years. A common model for accretion flows incorporating truncated thin disks is one where the thin disk transitions to an inner RIAF as the flow crosses a certain transition radius $r_{\\rm t}$ (Esin et al.\\ 1997); a commonly invoked type of RIAF is an advected-dominated accretion flow (ADAF), wherein the disk is optically thin and geometrically thick (e.g., Narayan \\& Yi 1995). The largest width observed for the Fe K$\\alpha$ line thus could indicate $r_{\\rm t}$. In this model, $r_{\\rm t}$ is expected to increase, and more of the inner thin disk evaporates, as the accretion rate relative to Eddington, $\\dot{m}$, decreases in a given object. Supporting evidence for this comes from timing observations of black hole X-ray Binary systems during outburst decay: characteristic temporal frequencies in the power spectral density function (PSD), such as peaks of Lorentzian components and/or quasi-periodic oscillations, migrate towards lower temporal frequencies as $\\dot{m}$ decreases and the source luminosity fades, as the source evolves through the low/hard spectral state into quiescence (e.g., Axelsson et al.\\ 2005, Belloni et al.\\ 2005).  In addition, the temperature and flux of the soft, thermally emitted component have been seen to decrease with $\\dot{m}$ in many sources (e.g., Gierli\\'{n}ski, Done \\& Page 2008).  However, the predicted relationship between $\\dot{m}$ and $r_{\\rm t}$ remains unclear. Yuan \\& Narayan (2004) empirically derive that compact sources accreting at $\\dot{m}$ near $10^{\\sim-2}$,  $10^{\\sim-4}$, and  $10^{\\sim-(6-7)}$ may be associated with values of $r_{\\rm t}$ near  $10^{\\sim(1-2)}$, $10^{\\sim(2-3)}$,  and $10^{\\sim(4-5)}$    $R_{\\rm g}$, respectively. Assuming a 2--10 keV flux in 2002 of $3.9 \\times 10^{-11}$ erg cm$^{-2}$ s$^{-1}$ from {\\it RXTE-PCA} monitoring, a luminosity distance of 41.3 Mpc (following Mould et al.\\ 2000, and using $H_{\\rm o}$ = 70 km s$^{-1}$ Mpc$^{-1}$ and $\\Lambda_{\\rm o}$ = 0.73), the 2--10 keV luminosity $L_{2-10} = 7 \\times 10^{42}$ erg s$^{-1}$. From Marconi et al.\\ (2004), an AGN with this $L_{2-10}$ has a bolometric luminosity $L_{\\rm bol} = 15 L_{2-10} = 1.1 \\times 10^{44}$ erg s$^{-1}$. The accretion relative to Eddington $\\dot{m}$ = $L_{\\rm bol}/L_{\\rm Edd}$ is thus estimated to be 0.15 for the 2002 {\\it XMM-Newton} observation. $\\dot{m}$ during the {\\it Suzaku} observation is thus 0.04. Meanwhile, Lu \\& Wang (2000) have derived $\\dot{m} \\sim 5.5\\%$ from SED fitting. These values of $\\dot{m}$ and our derived value of $r_{\\rm t}$ are not immediately consistent with the rough relation of Yuan \\& Narayan (2004), thus pointing toward models incorporating smaller values of $r_{\\rm t}$ (see below). Furthermore, most low-$\\dot{m}$ sources are radio loud, but NGC 4593 is not strongly radio-loud. Its 5 GHz flux has been measured to near $\\sim$2 mJy (e.g., Schmitt et al. 2001), and its B-band flux is $\\sim$6--16 mJy (e.g., McAlary et al.\\ 1983), so the radio loudness parameter, defined as the ratio of these two values, is $\\lesssim$3. Values $\\geq$10 define a source as radio-loud  (Kellermann et al.\\ 1989). A connection between $\\dot{m}$ (proportional to the observed X-ray flux) and $r_{\\rm t}$ in NGC 4593 is thus qualitative only as well as speculative, especially since we have only two model-dependent estimates of $r_{\\rm t}$.  There is also the question of whether the inner portions of a thin disk in AGN can evaporate and/or become radiatively inefficient on timescales of only a few years. As the accretion disks of BH XRBs are thought to evolve on timescales of at least hours to days, the corresponding timescales in NGC 4593 (black hole mass a factor of $10^{\\gtrsim5}$ higher) would be decades to centuries. On the other hand, Marscher et al.\\ (2002) interpreted rapid (duration of $\\sim$ a couple weeks) dips in the X-ray light curve of the radio-loud Seyfert 3C~120 as periods when the inner portion of the disk evaporated, each event leading to ejection of material into the relativistic jet and a corresponding radio flare about a month later.     The total Fe K$\\alpha$ line intensity decreased from $4.79 \\pm 0.48 \\times 10^{-5}$ (2002) to $2.93 \\pm 0.23 \\times 10^{-5}$ (2007) ph cm$^{-2}$ s$^{-1}$, a factor of 1.7, while the observed 4--10 keV flux decreased from $2.47\\times 10^{-11}$ (2002) to $0.69 \\times 10^{-11}$ (2007) erg cm$^{-2}$ s$^{-1}$, a factor of 3.6, i.e., roughly half of the total line flux has responded to continuum decrease. Modeling the Fe K$\\alpha$ line with a dual-Gaussian model attempted to separate the variable and non-variable emission components; in this model, a non-variable, narrow component is detected in both observations, and dominates the total line flux in the {\\it Suzaku} spectrum, while a broader component is detected only in the {\\it XMM-Newton} spectrum.  The best-fit width $\\sigma$ of the broad line was $177^{+84}_{-52}$ eV, corresponding to FWHM velocity of $19100^{+9100}_{-5600}$ km s$^{-1}$, implying $r = 3.2^{+3.3}_{-1.7} \\times 10^{12}$ m = $0.12^{+0.12}_{-0.06}$ lt-days, or $330^{+330}_{-180} R_{\\rm g}$. This estimate is inconsistent with B07's estimate of $\\sim$2000 $R_{\\rm g}$ based on the invariance of the Fe K$\\alpha$ line during the {\\it XMM-Newton} observation, but it is still consistent with the presence of a truncated thin disk ($r_{\\rm t} > 150 R_{\\rm g}$). $r_{\\rm t}$ could be invariant from 2002 to 2007, but an annular region on the thin disk spanning from $\\sim 300 R_{\\rm g}$ to 1000--5000 $R_{\\rm g}$ (outer radius obviously speculative) could be responding to the drop in illuminating continuum flux. If the drop in \\ion{Fe}{26} flux is real, then that line could also originate in this region. However, it is not clear in this model why material at $\\gtrsim 6000 R_{\\rm g}$ (yielding the narrow line component), contributing roughly half of the total line intensity in 2002, has not responded to the drop in continuum flux, as it is well within a week's light-travel time. The width of the narrow line had been fixed at $\\sigma=41$ eV in our modeling of the {\\it XMM-Newton} profile, but contributions from a more narrow component likely cannot be ruled out. Such distant material could be responding to a previous higher continuum flux. The 2--10 keV {\\it RXTE}-PCA monitoring light curve in fact showed a higher, more average flux level until roughly 300 days before the {\\it Suzaku} observation (Figure 7).  A final possibility that does not require evolution in $r_{\\rm t}$ is that the inner disk may have become be too ionized to transmit an Fe line. In the context of models with a hot, ionized skin (Nayakshin, Kazanas \\& Kallman 2000), a disk illuminated by a power-law continuum with a photon index near 1.6, similar to that in the {\\it Suzaku} observation, yields extremely weak Fe line emission.   Of course, both profile models are likely oversimplifications. The community could thus benefit from an X-ray observatory with $\\sim$ few eV resolution combined with a large effective area near 6 keV to resolve the various components of the Fe K$\\alpha$ core as a function of time and/or continuum flux level, if multiple components do indeed exist, as well as resolve the \\ion{Fe}{26} line.     \\subsection{The Compton Reflection Component}    %  A Compton shoulder was not significantly detected in either the {\\it Suzaku} or {\\it XMM-Newton} spectra; we find upper limits to Compton shoulder emission (first-scattering) of $\\sim$23$\\%$ of the core. It is thus not obvious from this limit alone whether bulk of the Fe K$\\alpha$ line originates in Compton-thick material, especially the degree to which the strength of the Compton shoulder depends on the geometry of the material. However, no relativistically broadened Fe K$\\alpha$ line has been confirmed in NGC 4593, but we confirm from the broadband {\\it Suzaku} spectrum the presence of a Compton reflection hump which thus must correspond to (at least some fraction of) the Fe K$\\alpha$ core emission. We can investigate if the measured strength $R$ of the Compton reflection hump, $1.08 \\pm 0.20$ (statistical uncertainty only; $\\pm$ 0.35 including the systematic uncertainty), can correspond to the observed Fe K$\\alpha$ line $EW$ of 255$\\pm$19 eV. Following George \\& Fabian (1991), one expects an $EW$ of 135 eV (using the abundances of Lodders 2003) to correspond to $R$ = 1 for an semi-infinite optically thick slab illuminated by a power-law continuum with $\\Gamma$ = 1.7 and assuming solar abundances and an inclination angle of 30$\\degr$ relative to the observer's line of sight. The observed values of $R$ and $EW$ are thus consistent if the Fe abundance relative to solar, $Z_{\\rm Fe}$, is about 1.7, which is not unreasonable. For a truncated disk, this could be explained by having the Comptonizing corona consist of numerous flares lying in a sandwich-like geometry just above/below the thin disk (e.g., Haardt et al.\\ 1994), such that the disk spans 2$\\pi$ sr of the sky as seen by each X-ray flare. However, very distant (pc-scale), Compton-thick material lying out of the line of sight, which cannot be ruled out as contributing to the observed Fe K emission profile, could also contribute to the total Compton reflection strength.  However, if the thin disk is truncated, then a semi-infinite slab may not be an appropriate geometry, especially if the central X-ray source is not immediately close to the reflecting disk. The $EW$ of a truncated disk will be lower, but will depend on the location of the illuminating X-ray source. If we assume that the illuminating X-ray source is located on the disk symmetry axis a height $h$ above the disk, we can use the $EW$ to constrain $h$. George \\& Fabian (1991, their Fig.\\ 15) demonstrate that the reflected flux is dominated by the region of the disk with $r/h \\sim$ 1--2. For a truncated thin disk with $r_{\\rm t}$ = several thousand $R_{\\rm g}$, $h$ must also be $\\sim$ several thousand $R_{\\rm g}$ above the black hole. A configuration in which the X-ray corona is associated with the base of a jet along the symmetry axis may thus be applicable, e.g., Markoff, Nowak \\& Wilms (2005). NGC 4593, like many Seyferts, is known to host a pc-scale radio component (size of $<$15 pc; Schmitt et al.\\ 2001).               \\subsection{Spectral Variability of the Broadband Components}    The primary-law component in Seyfert X-ray spectra is usually attributed to inverse Comptonization of soft seed photons. In an ADAF flow, thermal Comptonization is expected to dominate the X-ray emission unless the accretion rate is extremely low, in which case thermal bremsstrahlung emission dominates the X-ray spectrum (Narayan et al.\\ 1998, Narayan 2005). One of our main results is that while the primary power-law component has dropped in flux by a factor of almost 4, the soft excess has dropped in flux by a factor of $\\sim$20 between 2002 and 2007, ruling out an origin for the soft excess with a size greater than 5 lt-years. This drop is likely linked to the decrease in the primary X-ray power-law, i.e., it may be either a cause of an effect of it.    We explored two phenomenological models for the soft excess, bremsstrahlung and thermal Comptonization. In the presence of an ADAF flow, one can expect thermal bremsstrahlung emission with a temperature of $10^{9-11}$ K (Narayan \\& Yi 1995), but such temperatures are higher by over 2 orders of magnitude compared to the temperature derived from our model fits and by B07. We also modeled the soft excess as inverse Comptonization of soft seed photons with an assumed input temperature of 50 eV by an optically thin corona with a temperature $k_{\\rm B}T \\sim 50$ keV, again obtaining similar results to B07. The location of such a process is not clear, though it could occur in the ionized skin of the thin disk (e.g., Magdziarz et al.\\ 1998, Janiuk et al.\\ 2001), or at the base of an outflowing jet. If both the soft excess and hard X-ray power-law components originate via Comptonization of disk seed photons, a decrease in the intensity of those soft seed photons (e.g., from an increase in $r_{\\rm t}$) between 2002 and 2007 could yield a correlated drop in both component's flux. Another possibility is that the optical depth of the Comptonizing components may have changed."
    },
    "0910/0910.4359_arXiv.txt": {
        "abstract": "% We present results of the investigation of the nature of double periodic variables (DPVs). We have selected a sample of Galactic eclipsing DPVs for a multiwavelength photometric study aimed to reveal their nature. The short orbital periodicity and the cyclic variability are decoupled and separately investigated. Shapes of orbital light curves are consistent with semi-detached binaries. The amplitude of the long cycle is always larger in redder bandpasses. ",
        "introduction": "Double Periodic Variables (DPVs) is a group of binary stars with orbital period between 1 and 16~d, characterized by additional long cyclic variability in the range of 50-600~d, correlated with the orbital period (Mennickent et al.~2003, Mennickent \\& Ko{\\l}aczkowski 2009a). These stars were discovered during the search for Be stars in the Magellanic Clouds based on the OGLE variable stars catalogue. A careful survey of the OGLE and MACHO photometry, allowed to detect this kind of variability in about a hundred stars in the Magellanic Clouds. The first results obtained for one LMC DPV star, OGLE05155332-6925581, were recently published (Mennickent et al.~2008).  Investigating the ASAS photometric data we have found eleven Galactic DPVs. The light curves of four Galactic DPVs are shown in Fig.1. After the discovery of this new group of variable stars, we have initiated a multiwavelength observing campaign. ",
        "conclusions": "The analysis of the available data, suggests that DPVs are binaries in a mass exchange evolution stage. Probably the primary is rotating at the critical velocity and matter cumulates in a circumprimary disc. Disc size increases until certain instability ejects mass outside the binary system. As a result, in DPVs we observe the long-term variability (Mennickent \\& Ko{\\l}aczkowski 2009abc)."
    },
    "0910/0910.3786_arXiv.txt": {
        "abstract": "{ It is usually assumed that the ellipticity power spectrum measured in weak lensing observations can be expressed as an integral over the underlying matter power spectrum. This is true at order ${\\mathcal O}(\\Phi^2)$ in the gravitational potential.  We extend the standard calculation, constructing all corrections to order ${\\mathcal O}(\\Phi^4)$. There are four types of corrections: corrections to the lensing shear due to multiple-deflections; corrections due to the fact that shape distortions probe the reduced shear $\\gamma/(1-\\kappa)$ rather than the shear itself; corrections associated with the non-linear conversion of reduced shear to mean ellipticity; and corrections due to the fact that observational galaxy selection and shear measurement is based on galaxy brightnesses and sizes which have been (de)magnified by lensing.  We show how the previously considered corrections to the shear power spectrum correspond to terms in our analysis, and highlight new terms that were not previously identified.  All correction terms are given explicitly as integrals over the matter power spectrum, bispectrum, and trispectrum, and are numerically evaluated for the case of sources at $z=1$.  We find agreement with previous works for the ${\\mathcal O}(\\Phi^3)$ terms.  We find that for ambitious future surveys, the ${\\mathcal O}(\\Phi^4)$ terms affect the power spectrum at the $\\sim 1-5\\sigma$ level; they will thus need to be accounted for, but are unlikely to represent a serious difficulty for weak lensing as a cosmological probe. } ",
        "introduction": "Cosmic shear, the distortion of light from distant galaxies by the tidal gravitational field of the intervening large scale structure, is an excellent tool to probe the matter distribution in the universe. The statistics of the image distortions are related to the statistical properties of the large scale matter distribution and can thereby be used to constrain cosmology. Current results already demonstrate the power of cosmic shear observations at constraining the clustering amplitude $\\sigma_8$ and the matter density $\\Omega_{\\mr m}$ \\citep[e.g.,][]{Cosmos,Tim,CFHTLS, Fu}. Furthermore, cosmic shear provides an ideal tool to study dark energy through measuring the evolution of non-linear structure with large future surveys (DES\\footnote{www.darkenergysurvey.org}, LSST\\footnote{www.lsst.org},JDEM\\footnote{http://jdem.gsfc.nasa.gov}, Euclid\\footnote{http://sci.esa.int/science-e/www/object/index.cfm?fobjectid=42266}). These upcoming large weak lensing experiments will limit the statistical uncertainties to the percent level.  In order to extract cosmological information from these cosmic shear experiments, the increased data quality needs to be accompanied by a thorough treatment of systematic errors. On the observational side, this requires accurate information on the redshift distribution of source galaxies \\citep{MHH_photoz} and precise measurements of galaxy shapes which correct for observational systematics such as pixelization, noise, blurring by seeing and a spatially variable point spread function \\citep[see][]{STEP2,GREAT08}. On the theoretical side, astrophysical contaminants, like source lens clustering \\citep{BvWM97,Schneider02}, intrinsic alignment \\citep{King_ic} and the correlation between the gravitational shear and intrinsic ellipticities of galaxies \\citep{HS04, King_il,JS08,Zhang08,JS09}, need to be understood and removed. The prediction of lensing observables also requires precise models of the non-linear matter power spectrum and models for the relation between lensing distortion and large scale matter distribution which go beyond linear theory. While N-body simulations may predict the non-linear dark matter power spectrum with percent level accuracy in the near future \\citep{Coyote1, Coyote2}, the effect of baryons, which is a significant contamination to the weak lensing signal above $l\\sim 2000$ \\citep{Jing06, Rudd08}, is more difficult to account for and is the subject of ongoing work.  In this paper, we consider corrections to the relation between the observed lensing power spectra and the non-linear matter density field. In the regime of weak lensing, the observed galaxy ellipticities ($e_I$) are an estimator of the reduced shear $g_I = \\gamma_I/(1-\\kappa)$, \\be \\ensav{e_I}  = C\\frac{\\gamma_I}{1-\\kappa}\\; , \\label{eq:e} \\ee where $C$ is a constant which depends on the type of ellipticity estimator \\citep[e.g.][]{SS95, SS97} and the properties of the galaxy population under consideration, $\\gamma_I$ is a component of the shear, $\\kappa$ is the convergence, and the subscript $I$ refers to the two components of the ellipticity/shear \\citep[see e.g. ][for more details]{BS01}. The two-point statistics of the measured ellipticities are simply related to the reduced shear power spectrum. \\citet{CH} have calculated the shear power spectrum to fourth order in the gravitational potential. For the reduced shear power spectrum there exists an approximation to third order in the gravitational potential \\citep{DSW}. \\citet{Shapiro08} has demonstrated that on angular scales relevant for dark energy parameter estimates the difference between shear and reduced shear power spectra is at the percent level and ignoring these corrections will noticeably bias dark energy parameters inferred from future weak lensing surveys.  \\citet{FS_lens} introduced another type of corrections, termed \\emph{lensing bias}, which has a comparable effect on the shear power spectrum as the reduced shear correction: Observationally, shear is only estimated from those galaxies which are bright enough and large enough to be identified and to measure their shape. This introduces cuts based on observed brightness and observed size, both of which are (de)magnified by lensing \\citep[e.g.][]{BTP_mag, Jain_size}, and will thus bias the sampling of the cosmic shear field.  In the following we complete the calculation of the reduced shear power spectrum to fourth order in the gravitational potential to include multiple deflections and to account for the effects of lensing bias and the non-linear conversion between ellipticity and reduced shear.  We consider all lensing-related effects through ${\\mathcal O}(\\Phi^4)$, but do {\\em not} include effects associated with the sources (source clustering and intrinsic alignment corrections).  This paper is organized as follows: We describe our technique for calculating higher order lensing distortions and power spectra in Sect.~\\ref{sec:method}. Derivations of the different types of corrections to the shear and reduced shear power spectra are given in Sect.~\\ref{sec:shear} through Sect.~\\ref{sec:lensbias}. We quantify the impact of these corrections on future surveys in Sect.~\\ref{sec:impact} and discuss our results in Sect.~\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  We have calculated the reduced shear power spectra perturbatively to fourth order in the gravitational potential, accounting for the differences between shear and  reduced shear, relaxing the Born approximation, and including lens-lens coupling in the calculation of shear and convergence. The full set of corrections to the reduced shear power spectra are given in Table \\ref{tab:cg} (E-mode) and Table \\ref{tab:cgB} (B-mode).  The ellipticity power spectrum contains additional contributions, Eq.~(\\ref{eq:NLcorr}), which arises from the non-linearity of the shear estimator and depends on the specific definition of ellipticity used, and Eq.~(\\ref{eq:LBE}) which is caused by lensing bias. Through order $\\Phi^4$, this is the full set of corrections to the power spectrum arising from the lensing process itself. All corrections have been derived within the Limber approximation, and the analysis of ``12\" type multiple-deflection corrections is left for future work. Other corrections associated with the source galaxy population, such as source clustering and intrinsic alignments, are not treated in this paper. We find that, depending on the properties of the source galaxy population and on the type of shear estimator used, these corrections will be at the $\\sim 1-5\\sigma$ level, and thus should be included in the analysis of future precision cosmology weak lensing experiments.  That said, we caution that there are other areas in which the theory of weak lensing needs work if it is to meet ambitious future goals. Current fitting formula of the non-linear dark matter power spectrum have an accuracy of about $10\\%$ at arcminute scales \\citep{sm03} and the uncertainty exceeds $30\\%$ for $l>10000$ \\citep{HH09}, due to this difficulty in modeling the non-linear gravitational clustering angular scales of $l> 3000$ are likely to be excluded from parameter fits to cosmic shear measurements. Utilizing near-future $N$-body simulations it will become possible to determine the non-linear dark matter power spectrum with percent level accuracy \\citep[e.g.,][]{Coyote1, Coyote2}. However, this does not account for the effect of baryons, which will likely be important at halo scales and depend critically on the details of baryonic processes (cooling, feedback) involved. Baryons in dark matter halos which are able to cool modify the structure of the dark matter halo through adiabatic contraction \\citep{Blumenthal86,Gnedin04}, causing deviations of the inner halo profile from the simple NFW form and changing the halo mass - halo concentration relation \\citep[e.g.][]{Rudd08, Pedrosa09}. The latter can be constrained though galaxy-galaxy lensing \\citep{MandelbaumGG}, or could be internally self-calibrated in a weak lensing survey via its preferential effect on the small-scale power spectrum \\citep{Zentner08}. Baryons in the intergalactic medium may make up  about $10$\\% of the mass in the universe, and if their distribution on Mpc scales has been strongly affected by non-gravitational processes then they could pose a problem for precise calculation of the matter power spectrum \\citep[see][for an extreme and probably unrealistic example]{LG06}.  Given these uncertainties in modeling the non-linear matter distribution and that all the corrections derived in this paper are integrals over the non-linear matter power spectrum, bispectrum and trispectrum, we refrain from calculating $\\mathcal O(\\Phi^5)$ and higher corrections. We expect that the corrections derived in this paper are sufficient to model the perturbative relation between the non-linear matter distribution and the lensing distortion in weak lensing surveys for the forseeable future.  \\begin{acknowledgement} E.K. and C.H. are supported by the US National Science Foundation under AST-0807337 and the US Department of Energy under DE-FG03-02-ER40701. C.H. is supported by the Alfred P. Sloan Foundation. We thank Wayne Hu, Fabian Schmidt, Peter Schneider and Chaz Shapiro for useful discussions. \\end{acknowledgement}"
    },
    "0910/0910.1260_arXiv.txt": {
        "abstract": "We report new estimates of the time delays in the quadruple gravitationally lensed quasar PG1115+080, obtained from the monitoring data in filter $R$ with the 1.5-m telescope at the Maidanak Mountain (Uzbekistan, Central Asia) in 2004-2006. The time delays are 16.4 days between images C and B, and 12 days between C and A1+A2, with image C being leading for both pairs. The only known estimates of the time delays in PG1115 are those based on observations by Schechter et al. (1997) -- 23.7 and 9.4 days between images C and B, C and A1+A2, respectively, as calculated by Schechter et al., and 25 and 13.3 days as revised by Barkana (1997) for the same image components with the use of another method. The new values of time delays in PG 1115+080 may be expected to provide larger estimates of the Hubble constant thus decreasing a diversity between the $H_0$ estimates taken from gravitationally lensed quasars and with other methods. ",
        "introduction": "As was first suggested by Refsdal (1964), gravitationally lensed quasars can potentially provide an estimate of the Hubble constant $H_0$ independent of any intermediate distance ladder. This can be made from measurements of the time delays between the quasar intrinsic brightness variations seen in different quasar images. The value of $H_0$ can be obtained then (within the adopted cosmological model) from the observed geometry of the system, with the known lens and source redshifts, and with the use of physically validated model of mass distribution in the lensing galaxy. Since a phenomenon of gravitational lensing is controlled by the surface density of the total matter (dark plus luminous), it provides a unique possibility both to determine the value of $H_0$ and to probe the dark matter abundance in lensing galaxies and along the light paths in the medium between the quasar and observer.  By now the time delays have been measured in about 20 gravitationally lensed quasars resulting in the values of $H_0$, different for different quasars, while remaining noticeably less than the most recent estimate of $H_0$ obtained in the HST Hubble Constant Key Project with the use of Cepheids - $H_0=72\\pm8$ km s$^{-1}$ Mpc$^{-1}$ (Freedman et al. 2001). This discrepancy is large enough and, if the Hubble Constant is really a universal constant, needs to be explained. A detailed analysis of the problem of divergent $H_0$ estimates inherent in the time delay method and the ways to solve it can be found, e.g., in Keeton \\& Kochanek (1997); Kochanek (2002); Kochanek \\& Schechter (2004); Schechter (2005).  The quadruply lensed quasars, and PG1115+080 in particular, are known to better suit for determining the $H_0$ value as compared to the two-image lenses since they provide more observational constraints to fit the lens model. The PG1115 source quasar with a redshift of $z_S=1.722$ is lensed by a galaxy with $z_G=0.31$ (Henry \\& Heasley 1986; Christian et al. 1987; Tonry 1998), which forms four quasar images, with an image pair A1 and A2 bracketing the critical curve very close to each other. It is the second gravitationally lensed quasar discovered over a quarter of century ago, at first as a tripple quasar (Weymann et al. 1980). Hege et al. (1980) were the first to resolve the brightest image component into two images separated by 0.48 arcsec. Further observations (Young et al. 1981; Vanderriest et al. 1986; Christian et al. 1987; Kristian et al. 1993) have provided positions of quasar images and information about the lensing object, which  allowed to build a model of the system (e.g., Keeton, Kochanek \\& Seljak 1997). In particular, Keeton \\& Kochanek (1997) have shown that the observed quasar image positions and fluxes and the galaxy position can be fit well by an ellipsoidal galaxy with an external shear rather than by a just ellipsoidal galaxy or a circular galaxy with an external shear. They noted that a group of nearby galaxies detected by Young et al. (1981) could provide the needed external shear.  The problem of determining the Hubble constant from the time delay lenses is known to suffer from the so-called central concentration degeneracy, which means that, given the measured time delay values, the estimates of the Hubble constant turn out to be strongly model-dependent. In particular, models with more centrally concentrated mass distribution (lower dark matter content) provide higher values of $H_0$, more consistent with the results of the local $H_0$ measurements than those with lower mass concentration towards the centre (more dark matter).  \\begin{figure*} \\resizebox{\\hsize}{!}{\\includegraphics{Fig1.eps}} \\caption{The light curves of PG 1115+080 A1, A2, B, C from observations in filter $R$ with the 1.5-m telescope of the Maidanak Observatory in 2004, 2005 and 2006.} \\end{figure*}  The time delays in PG 1115+080 were determined for the first time by Schechter et al. (1997) to be $23.7\\pm 3.4$ days between B and C, and $9.4\\pm 3.4$ days between A1+A2 and C (image C is leading). Barkana (1997) re-analyzed their data using another algorithm and reported $25^{+3.3}_{-3.8}$ days for the time delay between B and C. There are also estimates of the time delays between images A1 and A2 made from the {\\it Chandra} and {\\it XMM-Newton} X-ray Observatories data (Dai X, et al. 2001, Chartas et al. 2004). As is predicted by all lens models, it does not exceed a small fraction of the day, and equals $0.16\\pm0.02$ days (Chandra data) or $0.149\\pm0.006$ days (XMM-Newton data). Determination of the time delays has generated a flow of models for the system, (Schechter et al. 1997; Keeton \\& Kochanek 1997; Courbin et al. 1997; Impey et al. 1998; Saha \\& Wiiliams 2001; Kochanek, Keeton \\& McLeod 2001; Zhao \\& Pronk 2001; Chiba 2002; Treu \\& Koopmans 2002; Yoo et al. 2005, 2006; Pooley et al. 2006; Miranda \\& Jetzer 2007), all illustrating how strongly the estimated value of $H_0$ depends on the adopted mass profiles of the lens galaxy for the given values of time delays. ",
        "conclusions": ""
    },
    "0910/0910.5597_arXiv.txt": {
        "abstract": "We apply the Standardized Candle Method (SCM) for Type II Plateau supernovae (SNe II-P), which relates the velocity of the ejecta of a SN to its luminosity during the plateau, to 15 SNe II-P discovered over the three season run of the Sloan Digital Sky Survey - II Supernova Survey.   The redshifts of these SNe - $0.027 < z < 0.144$ - cover a range hitherto sparsely sampled in the literature; in particular, our SNe II-P sample contains nearly as many SNe in the Hubble flow ($z > 0.01$) as all of the current literature on the SCM combined.   We find that the SDSS SNe have a very small intrinsic \\emph{I}-band dispersion ($0.22$ mag), which can be attributed to selection effects.  When the SCM is applied to the combined SDSS-plus-literature set of SNe II-P, the dispersion increases to $0.29$ mag, larger than the scatter for either set of SNe separately.  We show that the standardization cannot be further improved by eliminating SNe with positive plateau decline rates, as proposed in \\citet{P09}.  We thoroughly examine all potential systematic effects and conclude that for the SCM to be useful for cosmology, the methods currently used to determine the \\fe\\ velocity at day 50 must be improved, and spectral templates able to encompass the intrinsic variations of Type II-P SNe will be needed. ",
        "introduction": "Type Ia supernovae (SNe Ia) are standardized candles; although their peak magnitudes can vary by up to $\\sim 2.5$ mag, empirical relations between the shape of their light curves and peak magnitude result in distance measurements that can be accurate to $\\sim$7\\% \\citep{Phillips,JRK,Guy}.  This standardization has proven to be invaluable, as observations of SNe Ia led to the discovery of the accelerating expansion of the universe \\citep{Riess98,Perl99}.  Dedicated surveys have now observed over a thousand of these objects in the past several years over a wide range of redshifts, lowering statistical uncertainties with the volume of their discoveries (For high redshift SN programs, see \\citealp{Riess04,Riess07,astier06,WV07,miknaitis07,Kessler09,Frieman08}; for low redshift, see \\citealp{LOSS,CSP,SNF,CFA}).  As a result, selection effects and other potential sources of bias are of increasing importance to improving the constraints on the dark energy equation of state.  In order to fully utilize future surveys that will discover very large numbers of SNe (DES\\footnote{https://www.darkenergysurvey.org}; LSST\\footnote{http://www.lsst.org/lsst}), these systematic uncertainties will have to be reduced.   Type II Plateau Supernovae (SNe II-P) can also be used as distance indicators, though thus far with less precision and to much lower distances than SNe Ia.  SNe II-P are objects that, observationally, are classified as Type II SNe by the presence of hydrogen absorption in their spectra, and as `Plateau' based on the slow decay of their early light curves \\citep{barbon79}.  Distance measurements using the Expanding Photosphere Method \\citep[EPM; see][]{EPM1,EPM2,EPM3} and the more recent, related Spectral Expanding Atmosphere Method \\citep[SEAM; see][]{SEAM1,SEAM2}, which are based on the modeling of the SN atmosphere, have been shown to recover distances to 10\\% precision \\citep{SEAM3}.  A different approach has been taken by \\citet[hereafter HP02]{HP02}, who have shown that SNe II-P magnitudes can also be standardized empirically, despite being far more heterogeneous \\citep[peak luminosity dispersion of more than 5 mags; see][]{Patat,Pastorello} than SNe Ia.  The theoretically based methods take advantage of the relative simplicity of modeling hydrogen, which dominates the atmosphere of SNe II-P, as compared to the intermediate mass elements that make up the ejecta of SNe Ia.  SNe II-P also differ from their brighter brethren in a couple of other advantageous ways.  Over the past decade, observations in archival pre-SNe II-P images have found red supergiants in the mass range $~7-16\\, M_{\\odot}$ \\citep{IIP1,IIP2,IIP3,IIP4,IIP5,IIP6,IIP7,IIP8}, and recently it has been shown that in at least one of these cases, the supposed progenitor is no longer present after the SN II-P has dimmed \\citep{IIP9}.  So, compared to SNe Ia, where there is still uncertainty over the progenitor system (or systems), SNe II-P explosions are better understood.  Also, since SNe II-P have only been found in late-type galaxies - unlike SNe Ia, which are found in both late- and early-type galaxies - it is likely that biases from environmental effects will have a smaller effect on distance measurements for SNe II-P than SNe Ia.  Thus the differences between the two types of SNe and their methods of standardization will result in different systematic effects, allowing SNe II-P to serve as a useful check on SNe Ia measurements.   The standardization of the SNe II-P luminosity was introduced by HP02, who found that the luminosity and the expansion velocity at the photosphere are correlated when the SN is in its plateau phase; they use 50 rest-frame days post explosion as a convenient reference epoch to compare across SNe (this epoch is late enough to ensure the SN has entered the plateau phase, while still before it leaves the plateau).  The origin of this relationship is strongly grounded physically; a more luminous supernova's hydrogen recombination front will be at a greater distance from the core than in a less luminous SN, and thus the velocity of matter at the photosphere will be greater.  Using the velocity of the ejecta as measured from the minimum of the \\feII\\ absorption feature and a reddening correction based on the color of the SN when its light curve falls off the plateau, they found that the scatter in the Hubble diagram for \\emph{V} and \\emph{I} magnitudes drops from 0.95 and 0.80 mag, respectively, to 0.39 and 0.29 mag.  When they then restricted their sample to the 8 SNe with $z > 0.01$ in order to reduce the effect of peculiar velocities on host galaxy redshift, the scatter of their sample dropped even further, to 0.20 mag in \\emph{V} and 0.21 mag in \\emph{I}.  This technique has come to be known as the Standardized Candle Method (SCM).  \\citet[hereafter N06]{N06} improved upon the HP02 method in several ways that address its limitations for application to SNe at cosmological distances.  The HP02 host galaxy extinction correction is replaced by the rest-frame \\VmI\\ color at day 50, which can be obtained with less extensive late time photometric monitoring.  To allow spectroscopic follow-up programs to be more flexible, they also derive an observational relationship between the velocity of the \\fe\\ line at epochs spanning days 9-75 to the velocity at day 50, which is crucial for any realistic spectroscopic follow-up program.  They apply the standardization method to 5 high-redshift ($0.13 < z < 0.29$) SNe obtained as part of the Supernova Legacy Survey \\citep[SNLS;][]{astier06}, as well as one at $z = 0.019$ followed up by the Caltech Core-Collapse Program \\citep[CCCP;][]{CCCP} (SN 1999gi, which was added to the dataset of SNe II-P by \\citet{H03}, is also included).  The result of their SNLS-only fit differed strongly from that of the low-redshift SNe in HP02, which was attributed to a number of factors, including a small sample size, differences in data analysis techniques, and observational biases.  Their data did confirm, however, that the relation between ejecta velocity and magnitude could be used to reduce the intrinsic dispersion in the sample.  \\citet{Olivares} applied the SCM to 37 SNe II-P, for which they use 30 days before the end of the plateau as the common epoch of reference.  They find the Hubble Diagram dispersion to be remarkably similar whether one uses \\emph{B}-, \\emph{V}-, or \\emph{I}-band magnitudes (0.28, 0.31, and 0.32 mags, respectively), and raise the possibility that the magnitude-velocity relation may be quadratic instead of linear.  The most recent work on SNe II-P standardization was done by \\citet[hereafter P09]{P09}, which presented new data on 19 low-redshift (all located at $z < 0.03$) SNe II-P with multi-epoch spectra.  These SNe were discovered as part of the Lick Observatory SN Search \\citep[LOSS;][]{LOSS} by the Katzman Automatic Imaging Telescope (KAIT).  The most significant advancement of the SCM presented in P09 is the introduction of a robust technique for determining the velocity of the ejecta by cross-correlating the observed spectra with high-quality SNe II-P templates contained in SNID \\citep[SuperNova Identification Code; see][]{SNID1,SNID2}.  As mentioned in P09, the velocity of the weak, broad \\fe\\ line is difficult to measure without introducing systematic offsets that hinder comparisons across multiple data sets.  Using the SCM, they found a significantly larger \\emph{I}-band scatter (0.38 mag) in their Hubble Diagram than either HP02 or N06 (including all objects from both previous works), but showed that by removing all SNe with a positive plateau decline rate (defined as a decrease in brightness in \\emph{I}-band magnitude between days 10 and 50) from their analysis, the dispersion of this `culled' sample was only 0.22 mag, or $\\sim 10\\%$ in distance.   In this work we apply the SCM to both the 15 new SNe II-P from the SDSS-II Supernova Survey and the combined set of our SNe plus those in the literature.  Our sample fills in a redshift range, $0.027 < z < 0.144$, that has very few other objects.  We discuss our observations in \\S\\ref{sec-obs}, and present light curves for these SNe and an additional 19 spectroscopically confirmed SNe II-P, also discovered during the SDSS-II Supernova Survey, for which we do not have sufficient data to include in our SCM sample.  We also discuss how we determine which SNe are part of the SCM sample.  In Section \\S\\ref{sec-phot} we explain in detail our process for determining K- and S-corrections, and explore how these are affected by the uncertainty in the explosion date and the choice of spectral template.  We analyze the spectra of our SNe in \\S\\ref{sec-spec}, and obtain their ejecta velocity at day 50.  In \\S\\ref{sec-cosm} we present the results from combining our observed SNe with those of HP02, N06, and P09, and applying the SCM to a combined sample of 49 SNe II-P.  We discuss our results and their interpretations in \\S\\ref{sec-results}, and conclude with an eye towards what will be needed for future SNe II-P campaigns in \\S\\ref{sec-dis}. ",
        "conclusions": "\\label{sec-dis}  We have presented photometric and spectroscopic data for 15 SNe II-P from the SDSS-II Supernova Survey.  These SNe span a range of redshifts not covered by other surveys, connecting the local sample of HP02 and P09 to the high-$z$ SNLS sample in N06.  We propagate through our photometry reduction the uncertainty in the time of explosion in a way that accurately represents the state of knowledge regarding the explosion epoch.  We also perform K- and S-corrections for all of our SNe and describe the process in detail.  The best fit parameters for the SCM are updated with our new data, and are found to change significantly from those presented in P09.  The difference can be primarily attributed to the SDSS SNe being intrinsically brighter than the SNe in the `culled' sample of P09 (due to understood selection effects), without displaying the corresponding increase in ejecta velocity predicted by the SCM parameters of P09.  We show that the act of warping a SN II-P spectral template to the observed colors, while not completely removing systematic uncertainties, reduces them to such a level that they are unlikely to account for the magnitude of our offset.  The difficulty in accurately measuring the velocity of the broad \\fe\\ absorption features and the uncertainty involved in extrapolating velocities at early epochs to later times are likely to be the dominant sources of systematic uncertainty in our analysis.  The low $\\alpha$ value we derive for the SDSS-only sample is due to both the small scatter in observed magnitudes and the uncertainty in our velocity determination.  Despite the selection effects and measurement uncertainties that limit the power of the SDSS SNe II-P to constrain the SCM parameters, we observe the primary relationship defining the SCM, which is that brighter SNe have faster ejecta, at $\\approx 2\\sigma$.  The primary objective of the SDSS-II Supernova Survey was to identify and characterize SNe Ia; as a result, the observing strategy was not ideal for SNe II-P studies.  A future survey program that would focus on the SCM would greatly benefit from possessing the following qualities:  1)  Longer continuous survey duration.  For the final six weeks of such a survey, any newly discovered SNe will not have photometry sufficiently far into their plateau to allow extrapolation of their magnitudes at rest-frame day 50.  In addition, no SN can have its explosion date pinned down unless it happens after the first survey observation in its field.  Since each season of the SDSS-II SNS was three months long, no more than $50 \\%$ of the survey time was useful for discovering new SNe II-P.  Doubling the length of such a survey would nearly triple the number of SNe discovered for which the SCM can be used.  2)  Deeper, later spectroscopic observations.  When a new SN is discovered in a survey dedicated to SNe II-P, an initial quick spectrum is not needed.  A SN II-P will have a distinct light curve from a SN Ia, and thus valuable spectroscopic resources will not be used to follow up objects of no interest.  Since the velocity of the ejecta at day 50 is what is required, a spectrum at early times is not required, and observations can target day 50.  As we have shown, the systematic uncertainties associated with early-time spectra require a thorough study.  Also, each object must be observed for a sufficient time to obtain a high quality S/N spectrum.  Accurate measurements of \\feII\\ absorption lines are key to determining the strength of the $\\alpha$ parameter; quality over quantity should be the mantra for these followups.  3)  Template database must be extended.  To perform both K-corrections and S-corrections, an accurate representation of the true spectral energy distribution is vital.  Ideally, the spectrum for each SN in such a survey would have sufficient wavelength coverage to allow corrections for all bands of interest to be computed from it directly.  As this situation is unlikely, complementary spectral observations that cover the near IR should allow for templates to be constructed that range from at least $4000 - 12000$\\AA\\, (or redder still, if the wavelengths in the lowest energy filter are redshifted out of this range at the targeted $z$).  The amount of uncertainty in the day 50 magnitude results that is due to relying on a limited number of SNe II-P spectral templates will only be known when more extensive templates are available.  When these advances can be made, we will be able to determine whether the SCM and SNe II-P can place meaningful, independent constraints on cosmological parameters."
    },
    "0910/0910.5845_arXiv.txt": {
        "abstract": "We review three Li problems. First, the Li problem in the Sun, for which some previous studies have argued that it may be Li-poor compared to other Suns. Second, we discuss the Li problem in planet hosting stars, which are claimed to be Li-poor when compared to field stars. Third, we discuss the cosmological Li problem, i.e. the discrepancy between the Li abundance in metal-poor stars (Spite plateau stars) and the predictions from standard Big Bang Nucleosynthesis. In all three cases we find that the ``problems'' are naturally explained by non-standard mixing in stars. ",
        "introduction": "As illustrated in Fig. 1, the present day solar Li abundance \\citep[$A_{\\rm Li}$\\footnote{$A_X = \\log (N_X/N_H) + 12$} = 1.05;][]{asp09} is much lower than the meteoritic Li abundance \\citep[$A_{\\rm Li}$ = 3.26;][]{asp09}. This large depletion in the observed solar Li abundance by a factor of 160 relative to the primordial solar system composition remains one of the most serious challenges of standard solar models, which, as illustrated in Fig. 1, destroy only a minor amount (0.06 dex) of Li \\citep[e.g.][]{dan84}. It is important to note that modern models using updated OPAL opacities predict too much lithium destruction during the pre-main sequence \\citep[e.g.][]{pia02,dan03}, at variance with the observations of Li in solar-mass stars in young open clusters (see e.g. Fig. 6). When the new low solar abundances \\citep{asp09} are used the problem is reduced \\citep{ses06}, but in general classical models have problems dealing with pre-main sequence convection, so other parameters are invoked to reduce the efficient lithium destruction during the pre-main sequence \\citep[e.g.][]{ven98}.   \\subsection{Comparison to solar analogs}  A comparison between the Sun and solar analogs of one solar mass and solar metallicity \\citep{lam04} shows the Sun to be ``lithium-poor'' by a factor of 10. This apparent peculiarity in the solar Li abundance, led \\cite{lam04} to suggest that the Sun may be of dubious value for calibrating non-standard models of Li depletion. However, Pasquini et al. (1994) have shown that there are solar-type stars that have a Li abundance as low as in the Sun. Nevertheless, \\cite{pas94} comparison sample of solar-type stars span a wide range in stellar parameters (5400 K $<$ \\teff $<$ 6100 K, 3.6 $<$ log $g$ $<$ 4.6, -1.6 $<$ [Fe/H] $<$ +0.2) and therefore they are not representative of one-solar-mass solar analogs. In Fig. 2, we restrict the comparison sample of \\cite{pas94} to only solar analogs within $\\pm$200K in solar effective temperature, $\\pm$0.3 dex of the solar surface gravity and $\\pm$0.3 dex of the solar metallicity. As we can see, these solar analogs seem to cluster in two groups, one with very high lithium abundances of $A_{\\rm Li}$ $\\sim$ 2.4, i.e., 20 times higher than solar, and the other group with Li abundances as low as solar. Why are there no solar analogs with intermediate Li abundances? Why the number of solar analogs with low Li abundances is much lower than the number of analogs with high Li abundances? Could this be the reason why \\cite{lam04} only found stars with high Li abundances around one solar mass? The lack of stars with intermediate Li abundances in the Pasquini et al. (1994) solar analog sample, and the lack of stars with both intermediate and low Li abundance around one solar mass in the \\cite{lam04} sample are probably telling us that both samples have biases, perhaps due to a selection of stars mostly in one or two evolutionary stages, e.g., mainly young stars, which are known to have high Li abundances.  The recent work by \\cite{pas08} for solar analogs and solar twins in the solar-age open cluster M67, shows that solar twins (M67 stars around solar effective temperature) have Li abundance as low as solar, but stars 100 or 200 K hotter span a broad range in Li abundances. So, it is important that the temperature scale of the comparison sample is accurate; otherwise offsets of about 100 K may introduce a bias in the comparison between the Sun and stars.    \\subsection{Comparison to solar twins}  Solar twins, stars with stellar parameters very similar to the Sun, are ideal targets to see if the Sun is normal (or not) in its Li abundance. Being so similar to the Sun, it is possible to obtain reliable stellar parameters, and provided the Sun and the twins are analyzed (and observed) in a consistent way, the temperature scale is accurate. Furthermore, being selected due to their similarity in colors and luminosity to the Sun, they should span a range of ages very close to solar, avoiding thus potential biases in the selection of comparison stars in only one evolutionary stage very different to the present Sun.  Solar twins have been searched for a long time, and although interesting solar twin candidates like 16 Cyg B (HD 186427) were identified in the past, detailed analysis showed that they were significantly different to the Sun \\citep[see review by][]{cay96}.  When the first close solar twin (18 Sco) was found \\citep{por97}, it seemed to have a Li abundance near solar, but much better data \\citep[e.g.][]{mel06} showed that its Li abundance is actually three times higher than solar. One solar twin is certainly not an acceptable number for a comparison between the Sun and stars, so a large survey of solar twins was urgently needed. The two largest recent efforts for finding field solar twin stars are being undertaking by the group of Y. Takeda \\citep[e.g.][]{tak07} and by our group \\citep{mel06,mel07,mel09a,ram09}. Importantly, whenever possible, we are obtaining very high S/N for our sample stars, because otherwise only upper limits can be obtained for their Li abundances. Indeed, as shown by \\cite{tak07} in their Fig. 12, their data with S/N $\\sim$150 can only estimate upper limits for stars with $A_{\\rm Li}$ $<$ 1.5, i.e., they can only reliably determine Li abundances when they are three times higher than solar.  Our solar twin survey has been performed mainly with the 2.7m telescope at McDonald observatory in the North and with the 6.5m Magellan Clay telescope at Las Campanas observatory in the South. We have also obtained some Keck+HIRES data in the North and VLT+UVES and HARPS data in the South. Our data has been taken at R = 60,000-110,000 and achieving S/N = 200-1000. The first pilot data set taken at Keck resulted in the discovery of the second best solar twin, HD 98618 \\citep{mel06}, about a decade after the discovery of the first solar twin 18 Sco. HD 98618 seems to be a solar twin as good as 18 Sco, and, as this twin, it has also a Li abundance three times higher than solar. Learning from the experience of our pilot Keck observations, we improved our criteria to select the best solar twins, empirically adjusting our \\teff scale \\citep{ram05} for an apparent zero-point problem \\citep{cas09}. This is probably the reason why our first solar twin run at McDonald was very successful. Besides confirming the solar twin nature of 18 Sco and HD 98618, we identified two additional solar twins, HIP 56948 and HIP 73815 \\citep{mel07}, both with a low Li abundance similar to solar. HIP56948 remains to this date the star that most closely resembles the Sun, with a \\teff similar to solar within 10 K, as recently confirmed by \\cite{tak09} using Subaru+HDS observations. The year 2007 was very prolific for solar twin studies, besides the twins found by our group, \\cite{tak07} reported the discovery of the fifth solar twin, HIP 110963. Interestingly, this solar twin has a high Li abundance of $A_{\\rm Li}$ = 1.7. These five solar twins were starting to fill the Li desert seen in Fig. 2.  The year 2009 has been even better, with many more solar twins found using our McDonald data \\citep{ram09} and Magellan+MIKE observations \\citep{mel09a}, which together have found more than 30 stars very similar to the Sun. Thus, we are starting to have a large sample of solar twins for meaningful statistics, in particular for addressing the long-standing question on chemical peculiarities in the Sun \\citep{mel09a,ram09,gus09}.  In Fig. 3 we show the Li abundance for solar twins and solar analogs that have been analyzed for Li in our survey. Most stars from the Magellan run have already been analyzed, but the McDonald Li analysis is just starting. Open circles show detections and filled circles upper limits. The Sun is also shown for comparison. The first important point to note in this plot is that, unlike the lack of stars with both intermediate and low Li in the \\cite{lam04} sample, and the lack of stars with intermediate Li abundances in the \\cite{pas94} sample, our sample nicely covers a broad range of Li abundances 0.6 $<$ $A_{\\rm Li}$ $<$ 2.4, meaning probably that our sample is not affected by any significant selection bias.  In Fig. 4 we show the Li abundances vs. \\teff of our solar twin and solar analog sample restricted to stars with mass within $\\pm$3\\% solar and [Fe/H] within $\\pm$0.1 dex solar. The Sun does not look peculiar on this plot. One star with relatively high Li abundance stands out. As shown in Fig. 5, where Li is plotted as a function of age, the high Li abundance of this star is due to its young age. This plot has a smaller number of stars than Figs. 3-4 because here, besides the constraint to $\\pm$3\\% in mass and $\\pm$0.1 dex in [Fe/H], we show only stars for which we could determine reliable ages ($\\ge$ 2.5 sigma). This figure definitely shows that the solar Li abundance is not abnormal, at least not for a solar-metallicity solar-age one-solar-mass star, which at about 4.6 Gyr has already depleted a significant fraction of its original Li abundance.   A comparison of our solar twin results to one-solar-mass stars in solar metallicity ($\\pm$0.15 dex) open clusters \\citep[selected from the sample of][and including the results from Pasquini et al. 2008]{ses05} is shown in Fig. 6. Again, the agreement is excellent, and reinforces a strong correlation between Li depletion and age for one-solar-mass stars.  Non-standard models \\citep[e.g.][]{mon00,chat05,xio09,don09} can reproduce the observed data, as shown in Figs. 7 and 8. We are in the process of obtaining better Li abundances and ages for our sample stars, which can potentially constrain the range of initial rotational velocities of our solar twins.  Based on the results shown above, we conclude that the solar Li abundance is not peculiar but a product of depletion due to non-standard mixing which affect both the Sun and the solar twins. ",
        "conclusions": ""
    },
    "0910/0910.0717_arXiv.txt": {
        "abstract": "Although today there are many observational methods, Type Ia supernovae (SNIa) is still one of the most powerful tools to probe the mysterious dark energy (DE). The most recent SNIa datasets are the 307 SNIa ``Union'' dataset \\cite{kow08} and the 397 SNIa ``Constitution'' dataset \\cite{hic09}. In a recent work \\cite{wei10}, Wei pointed out that both Union and Constitution datasets are in tension with the observations of cosmic microwave background (CMB) and baryon acoustic oscillation (BAO), and suggested that two truncated versions of Union and Constitution datasets, namely ``UnionT'' and ``ConstitutionT'', should be used to constrain various DE models. But in \\cite{wei10}, only the $\\Lambda$CDM model is used to select the outliers from the Union and the Constitution dataset. In principle, since different DE models may select different outliers, the truncation procedure should be performed for each different DE model. In the present work, by performing the truncation procedure of \\cite{wei10} for 10 different models, we demonstrate that the impact of different models is negligible, and the approach adopted in \\cite{wei10} is valid. Moreover, by using the 4 SNIa datasets mentioned above, as well as the observations of CMB and BAO, we perform best-fit analysis on the 10 models. It is found that: (1) For each DE model, the truncated SNIa datasets not only greatly reduce $\\chi _{min}^{2}$ and $\\chi _{min}^{2}/dof$, but also remove the tension between SNIa data and other cosmological observations. (2) The CMB data is very helpful to break the degeneracy among different parameters, and plays a very important role in distinguishing different DE models. (3) The current observational data are still too limited to distinguish all DE models. ",
        "introduction": "Observations of Type Ia supernovae (SNIa) \\cite{rie98,per99,ton03,kno03,rie04a}, cosmic microwave background (CMB) \\cite{ben03,spe03,spe07,pag07,hin07,kom09} and large scale structure (LSS) \\cite{teg04a,teg04b,teg06} all indicate the existence of dark energy (DE) driving the current accelerating expansion of the universe. The most obvious theoretical candidate of DE is the cosmological constant $\\Lambda$, but it is plagued with the fine-tuning problem and the coincidence problem \\cite{wein89,wein00,sah00,car01,pee03,pad03,cop06}. There are also many dynamical DE models, such as quintessence \\cite{pee88,rat88,rat88,zla99}, phantom \\cite{cal02,car03}, $k$-essence \\cite{arm99,chi00,arm01}, CPL \\cite{che01,lin03}, tachyon \\cite{pad02,bag03}, hessence \\cite{wei05a,wei05b}, Chaplygin gas \\cite{kam01}, generalized Chaplygin gas \\cite{bto02}, holographic \\cite{li04,hua05a,hua05b,li08,li09a,zhax10}, agegraphic \\cite{wei08a,wei08b}, holographic Ricci \\cite{gao09}, Yang-Mills condensate \\cite{zhay07a,zhay07b,wans08a,wans08b}, etc. Although numerous theoretical models have been proposed in the past decade, the nature of DE still remains a mystery.  In recent years, the numerical study of DE, i.e. utilizing cosmological observations to constrain DE models, has become one of the most active fields in the modern cosmology \\cite{alb06,wany00,wany04,wany07,hut03,hut05,hua04,zhax04,zhax05,cha06,zhax07a,zhax07b,zhax09,wanb05,wanb06,ma09a,ma09b,wei07,wei08c,wei09}. Although today there are many observational methods, SNIa is still one of the most powerful tools to probe the mysterious DE. In the past decade, many SNIa datasets, such as Gold04 \\cite{rie04b}, Gold06 \\cite{rie07}, SNLS \\cite{ast07}, ESSENCE \\cite{woo07}, Davis \\cite{dav07}, have been released, while the number and quality of SNIa have continually increased. The most recent SNIa datasets are ``Union'' \\cite{kow08} and ``Constitution'' \\cite{hic09}, and they have been widely used in the literature \\cite{shaf09,li09b,qi09,hua09,wu09,chen09,gong10a,gong10b,san09,mor09,wans10}. However, these SNIa datasets are not always consistent with other types of cosmological observations, and are even in tension with other SNIa samples. For examples, in \\cite{jas05,jas06,ness05}, the Gold04 dataset was shown to be inconsistent with the SNLS dataset: the SNLS dataset favors the $\\Lambda$CDM model, while the Gold04 dataset favors the dynamical DE model. In \\cite{ness07}, by comparing the maximum likelihood fits of the CPL parameters ($w_0$, $w_1$) given by different SNIa samples, Nesseris and Perivolaropoulos found that the Gold06 dataset is also in $2\\sigma$ tension with the SNLS dataset. Moreover, they also investigated how to remove this tension. The method is simple. First, they fitted the $\\Lambda$CDM model to the whole 182 SNIa in the Gold06 dataset, and obtained the best-fit parameter value of $\\Lambda$CDM model (for SNIa data only). Then, they calculated the relative deviation to the best-fit $\\Lambda$CDM prediction, $|\\mu_{obs}-\\mu_{\\Lambda CDM}|/\\sigma_{obs}$, for all the 182 data points. Here $\\mu_{obs}$ is the observational value of distance modulus, $\\mu_{\\Lambda CDM}$ is the theoretical value of distance modulus given by the best-fit $\\Lambda$CDM model, and $\\sigma_{obs}$ is the 1$\\sigma$ error of distance modulus. By searching the SNIa samples satisfying $|\\mu_{obs}-\\mu_{\\Lambda CDM}|/\\sigma_{obs}>1.8$, they isolated six SNIa that are mostly responsible for the tension. Further, by using the random truncation method, they demonstrated that these 6 SNIa are systematically different from the Gold06 dataset.   In a recent work, by comparing the maximum likelihood fits of the CPL parameters ($w_0$, $w_1$), Wei \\cite{wei10} pointed out that both Union and Constitution dataset are also in tension with the observations of CMB and baryon acoustic oscillation (BAO). Moreover, he also investigated how to remove these tensions. By using the method of truncation of \\cite{ness07}, Wei found out the main sources that are responsible for the tensions: for the Union set, there are 21 SNIa differing from the best-fit $\\Lambda$CDM prediction beyond $1.9\\sigma$ (i.e. $|\\mu_{obs}-\\mu_{\\Lambda CDM}|/\\sigma_{obs}>1.9$); and for the Constitution set, there are 34 SNIa differing from the best-fit $\\Lambda$CDM prediction beyond $1.9\\sigma$. The specific limit of truncation (i.e. $1.9\\sigma$) is chosen based on two considerations: first, the tension between SNIa samples and other observations can be completely removed; second, the number of usable SNIa can be preserved as much as possible (see \\cite{wei10} for details). By subtracting these outliers from the Union dataset and the Constitution dataset, respectively, two new SNIa datasets, ``UnionT'' and ``ConstitutionT'' (``T'' stands for ``truncated''), were obtained. Further, Wei argued that the UnionT and the ConstitutionT datasets are fully consistent with the other cosmological observations, and should be used to constrain various DE models. But in \\cite{wei10}, only the $\\Lambda$CDM model is used to select the outliers from the Union and the Constitution dataset. Since different DE models may select different outliers, one may doubt whether the approach adopted in \\cite{wei10} is valid. In principle, the truncation procedure should be performed for each different DE model. Only if the impact of different models is negligible, one can conclude that the conclusion of Wei is correct. So in the present work, we shall consider 10 different models. The truncation procedure will be performed for all these 10 models, and the corresponding cosmological consequences will be explored.   This paper is organized as follows: In Section 2, we briefly describe 10 theoretical models considered in this work. In Section 3, we present the method of data analysis, as well as the SNIa datasets we used in this paper. By performing the truncation procedure of \\cite{wei10} for 10 different models, we demonstrate that the approach adopted in \\cite{wei10} is valid. In Section 4, we show the data fitting results of 10 models, and present the corresponding conclusions. Section 5 is a short summary. In this work, we assume today's scale factor $a_{0}=1$, so the redshift $z$ satisfies $z=a^{-1}-1$; the subscript ``0'' always indicates the present value of the corresponding quantity, and the unit with $c=\\hbar=1$ is used. ",
        "conclusions": "We will present the results of data fitting in this section. Using the Union, the UnionT, the Constitution and the ConstitutionT dataset, respectively, we list the $\\chi_{min}^{2}$ and the $\\chi_{min}^{2}/dof$ for those 10 models, in table \\ref{table1}, table \\ref{table2}, table \\ref{table3}, and table \\ref{table4}. Here ``SNIa'' means only SNIa data is used in the analysis, ``SNIa+BAO'' means both SNIa data and BAO data are used, ``SNIa+CMB'' means both SNIa data and CMN data are used, and ``SNIa+BAO+CMB'' means all these three types of observational data are taken into account. The main conclusions are summarized as follows.   \\begin{table} \\caption{The $\\chi_{min}^{2}$ and the $\\chi_{min}^{2}/dof$ (in the Parentheses) for the 10 models, where the Union dataset is used.} \\begin{center} \\label{table1} \\begin{tabular}{ccccc} \\hline\\hline ~~~Model~~~ & ~~~SNIa~~~ & ~~~SNIa+BAO~~~  & ~~~SNIa+CMB~~~ &  ~~~SNIa+BAO+CMB~~~ \\\\ \\hline\\hline ~~~$\\Lambda$CDM~~~ & ~~~$311.936~(1.019)$~~~ & ~~~$313.205~(1.020)$~~~ &  ~~~$312.424~(1.018)$~~~ &  ~~~$313.594~(1.018)$~~~ \\\\ \\hline ~~~DGP~~~ & ~~~$313.026~(1.023)$~~~ & ~~~$314.319~(1.024)$~~~ &  ~~~$339.280~(1.105)$~~~ &  ~~~$341.584~(1.109)$~~~ \\\\ \\hline ~~~ADE~~~ & ~~~$313.536~(1.025)$~~~ & ~~~$314.838~(1.026)$~~~ &  ~~~$327.216~(1.066)$~~~ &  ~~~$329.252~(1.069)$~~~ \\\\ \\hline ~~~XCDM~~~ & ~~~$310.682~(1.019)$~~~ & ~~~$311.966~(1.020)$~~~ &  ~~~$312.224~(1.020)$~~~ &  ~~~$313.456~(1.021)$~~~ \\\\ \\hline ~~~CG~~~ & ~~~$310.434~(1.018)$~~~ & ~~~$311.628~(1.018)$~~~ &  ~~~$311.921~(1.019)$~~~ &  ~~~$313.193~(1.020)$~~~ \\\\ \\hline ~~~HDE~~~ & ~~~$310.827~(1.019)$~~~ & ~~~$312.149~(1.020)$~~~ &  ~~~$311.218~(1.017)$~~~ &  ~~~$312.481~(1.018)$~~~ \\\\ \\hline ~~~RDE~~~ & ~~~$310.682~(1.019)$~~~ & ~~~$311.966~(1.020)$~~~ &  ~~~$312.323~(1.021)$~~~ &  ~~~$313.801~(1.022)$~~~ \\\\ \\hline ~~~LP~~~ & ~~~$309.984~(1.020)$~~~ & ~~~$310.744~(1.019)$~~~ &  ~~~$310.691~(1.019)$~~~ &  ~~~$311.978~(1.020)$~~~ \\\\ \\hline ~~~CPL~~~ & ~~~$310.091~(1.020)$~~~ & ~~~$310.896~(1.019)$~~~ &  ~~~$310.906~(1.019)$~~~ &  ~~~$312.258~(1.020)$~~~ \\\\ \\hline ~~~GCG~~~ & ~~~$310.405~(1.021)$~~~ & ~~~$311.401~(1.021)$~~~ &  ~~~$311.925~(1.023)$~~~ &  ~~~$313.325~(1.024)$~~~ \\\\ \\hline\\hline \\end{tabular} \\end{center} \\end{table}  \\begin{table} \\caption{The $\\chi_{min}^{2}$ and the $\\chi_{min}^{2}/dof$ (in the Parentheses) for the 10 models, where the UnionT dataset is used.} \\begin{center} \\label{table2} \\begin{tabular}{ccccc} \\hline\\hline ~~~Model~~~ & ~~~SNIa~~~ & ~~~SNIa+BAO~~~  & ~~~SNIa+CMB~~~ &  ~~~SNIa+BAO+CMB~~~ \\\\ \\hline\\hline ~~~$\\Lambda$CDM~~~ & ~~~$204.568~(0.718)$~~~ & ~~~$205.945~(0.720)$~~~ &  ~~~$205.575~(0.719)$~~~ &  ~~~$206.794~(0.721)$~~~ \\\\ \\hline ~~~DGP~~~ & ~~~$205.234~(0.720)$~~~ & ~~~$206.634~(0.722)$~~~ &  ~~~$226.987~(0.794)$~~~ &  ~~~$229.388~(0.799)$~~~ \\\\ \\hline ~~~ADE~~~ & ~~~$205.560~(0.721)$~~~ & ~~~$206.963~(0.724)$~~~ &  ~~~$216.201~(0.756)$~~~ &  ~~~$218.305~(0.761)$~~~ \\\\ \\hline ~~~XCDM~~~ & ~~~$204.058~(0.719)$~~~ & ~~~$205.449~(0.721)$~~~ &  ~~~$204.855~(0.719)$~~~ &  ~~~$206.207~(0.721)$~~~ \\\\ \\hline ~~~CG~~~ & ~~~$204.050~(0.718)$~~~ & ~~~$205.365~(0.721)$~~~ &  ~~~$204.552~(0.718)$~~~ &  ~~~$205.930~(0.720)$~~~ \\\\ \\hline ~~~HDE~~~ & ~~~$204.070~(0.719)$~~~ & ~~~$205.496~(0.721)$~~~ &  ~~~$204.228~(0.717)$~~~ &  ~~~$205.614~(0.719)$~~~ \\\\ \\hline ~~~RDE~~~ & ~~~$204.058~(0.719)$~~~ & ~~~$205.449~(0.721)$~~~ &  ~~~$205.875~(0.722)$~~~ &  ~~~$207.498~(0.726)$~~~ \\\\ \\hline ~~~LP~~~ & ~~~$204.055~(0.721)$~~~ & ~~~$205.347~(0.723)$~~~ &  ~~~$204.429~(0.720)$~~~ &  ~~~$205.739~(0.722)$~~~ \\\\ \\hline ~~~CPL~~~ & ~~~$204.057~(0.721)$~~~ & ~~~$205.391~(0.723)$~~~ &  ~~~$204.063~(0.719)$~~~ &  ~~~$205.517~(0.721)$~~~ \\\\ \\hline ~~~GCG~~~ & ~~~$204.053~(0.721)$~~~ & ~~~$205.366~(0.723)$~~~ &  ~~~$204.545~(0.720)$~~~ &  ~~~$206.090~(0.723)$~~~ \\\\ \\hline\\hline \\end{tabular} \\end{center} \\end{table}   \\begin{table} \\caption{The $\\chi_{min}^{2}$ and the $\\chi_{min}^{2}/dof$ (in the Parentheses) for the 10 models, where the Constitution dataset is used.} \\begin{center} \\label{table3} \\begin{tabular}{ccccc} \\hline\\hline ~~~Model~~~ & ~~~SNIa~~~ & ~~~SNIa+BAO~~~  & ~~~SNIa+CMB~~~ &  ~~~SNIa+BAO+CMB~~~ \\\\ \\hline\\hline ~~~$\\Lambda$CDM~~~ & ~~~$465.604~(1.176)$~~~ & ~~~$466.902~(1.176)$~~~ &  ~~~$466.316~(1.175)$~~~ &  ~~~$467.525~(1.175)$~~~ \\\\ \\hline ~~~DGP~~~ & ~~~$466.122~(1.177)$~~~ & ~~~$467.433~(1.177)$~~~ &  ~~~$498.264~(1.255)$~~~ &  ~~~$500.368~(1.257)$~~~ \\\\ \\hline ~~~ADE~~~ & ~~~$466.275~(1.177)$~~~ & ~~~$467.580~(1.178)$~~~ &  ~~~$483.675~(1.218)$~~~ &  ~~~$485.585~(1.220)$~~~ \\\\ \\hline ~~~XCDM~~~ & ~~~$465.602~(1.179)$~~~ & ~~~$466.901~(1.179)$~~~ &  ~~~$465.657~(1.176)$~~~ &  ~~~$466.947~(1.176)$~~~ \\\\ \\hline ~~~CG~~~ & ~~~$465.293~(1.178)$~~~ & ~~~$466.564~(1.178)$~~~ &  ~~~$465.591~(1.176)$~~~ &  ~~~$466.891~(1.176)$~~~ \\\\ \\hline ~~~HDE~~~ & ~~~$465.769~(1.179)$~~~ & ~~~$467.080~(1.179)$~~~ &  ~~~$466.030~(1.177)$~~~ &  ~~~$467.384~(1.177)$~~~ \\\\ \\hline ~~~RDE~~~ & ~~~$465.602~(1.179)$~~~ & ~~~$466.901~(1.179)$~~~ &  ~~~$472.841~(1.194)$~~~ &  ~~~$474.466~(1.195)$~~~ \\\\ \\hline ~~~LP~~~ & ~~~$461.071~(1.170)$~~~ & ~~~$461.570~(1.169)$~~~ &  ~~~$465.610~(1.179)$~~~ &  ~~~$466.910~(1.179)$~~~ \\\\ \\hline ~~~CPL~~~ & ~~~$461.526~(1.171)$~~~ & ~~~$462.112~(1.170)$~~~ &  ~~~$465.636~(1.179)$~~~ &  ~~~$466.902~(1.179)$~~~ \\\\ \\hline ~~~GCG~~~ & ~~~$465.080~(1.180)$~~~ & ~~~$466.360~(1.181)$~~~ &  ~~~$465.920~(1.180)$~~~ &  ~~~$466.895~(1.179)$~~~ \\\\ \\hline\\hline \\end{tabular} \\end{center} \\end{table}     \\begin{table} \\caption{The $\\chi_{min}^{2}$ and the $\\chi_{min}^{2}/dof$ (in the Parentheses) for the 10 models, where the ConstitutionT dataset is used.} \\begin{center} \\label{table4} \\begin{tabular}{ccccc} \\hline\\hline ~~~Model~~~ & ~~~SNIa~~~ & ~~~SNIa+BAO~~~  & ~~~SNIa+CMB~~~ &  ~~~SNIa+BAO+CMB~~~ \\\\ \\hline\\hline ~~~$\\Lambda$CDM~~~ & ~~~$269.081~(0.743)$~~~ & ~~~$270.610~(0.745)$~~~ &  ~~~$271.423~(0.748)$~~~ &  ~~~$272.764~(0.749)$~~~ \\\\ \\hline ~~~DGP~~~ & ~~~$269.443~(0.744)$~~~ & ~~~$270.979~(0.747)$~~~ &  ~~~$292.067~(0.805)$~~~ &  ~~~$294.360~(0.809)$~~~ \\\\ \\hline ~~~ADE~~~ & ~~~$269.613~(0.745)$~~~ & ~~~$271.139~(0.747)$~~~ &  ~~~$280.090~(0.772)$~~~ &  ~~~$282.147~(0.775)$~~~ \\\\ \\hline ~~~XCDM~~~ & ~~~$269.066~(0.745)$~~~ & ~~~$270.599~(0.748)$~~~ &  ~~~$269.245~(0.744)$~~~ &  ~~~$270.754~(0.746)$~~~ \\\\ \\hline ~~~CG~~~ & ~~~$268.992~(0.745)$~~~ & ~~~$270.502~(0.747)$~~~ &  ~~~$269.066~(0.743)$~~~ &  ~~~$270.594~(0.745)$~~~ \\\\ \\hline ~~~HDE~~~ & ~~~$269.113~(0.745)$~~~ & ~~~$270.664~(0.748)$~~~ &  ~~~$269.130~(0.743)$~~~ &  ~~~$270.698~(0.746)$~~~ \\\\ \\hline ~~~RDE~~~ & ~~~$269.066~(0.745)$~~~ & ~~~$270.599~(0.748)$~~~ &  ~~~$273.305~(0.755)$~~~ &  ~~~$275.180~(0.758)$~~~ \\\\ \\hline ~~~LP~~~ & ~~~$268.812~(0.747)$~~~ & ~~~$270.021~(0.748)$~~~ &  ~~~$269.068~(0.745)$~~~ &  ~~~$270.603~(0.748)$~~~ \\\\ \\hline ~~~CPL~~~ & ~~~$268.900~(0.747)$~~~ & ~~~$270.147~(0.748)$~~~ &  ~~~$269.122~(0.745)$~~~ &  ~~~$270.676~(0.748)$~~~ \\\\ \\hline ~~~GCG~~~ & ~~~$268.980~(0.747)$~~~ & ~~~$270.473~(0.749)$~~~ &  ~~~$269.069~(0.745)$~~~ &  ~~~$270.599~(0.748)$~~~ \\\\ \\hline\\hline \\end{tabular} \\end{center} \\end{table}  (1) For each DE model, the truncated SNIa datasets not only greatly reduce $\\chi _{min}^{2}$ and $\\chi _{min}^{2}/dof$, but also remove the tension between SNIa data and other cosmological observations.  Since it is too prolix to describe the results of all 10 models, in the following we will just discuss the HDE model as an example. From table \\ref{table1} and table \\ref{table2}, we can see that the combined Union+BAO+CMB data gives a $\\chi_{min}^{2}=312.481$ and a $\\chi_{min}^{2}/dof=1.018$, while the combined UnionT+BAO+CMB data gives a $\\chi_{min}^{2}=205.614$ and a $\\chi_{min}^{2}/dof=0.719$. This means that by dropping the 21 outliers, The $\\chi_{min}^{2}$ and the $\\chi_{min}^{2}/dof$ of the HDE model can reduce 106.867 and 0.299, respectively. From table \\ref{table3} and table \\ref{table4}, similar results are obtained. The combined Constitution+BAO+CMB data gives a $\\chi_{min}^{2}=467.384$ and a $\\chi_{min}^{2}/dof=1.177$, while the combined ConstitutionT+BAO+CMB data gives a $\\chi_{min}^{2}=270.698$ and a $\\chi_{min}^{2}/dof=0.746$. This means that by dropping the 34 outliers, The $\\chi_{min}^{2}$ and the $\\chi_{min}^{2}/dof$ of the HDE model can reduce 196.686 and 0.431, respectively. It should be pointed out that for all these 10 models, the UnionT and the ConstitutionT dataset can greatly reduce the corresponding $\\chi _{min}^{2}$ and $\\chi _{min}^{2}/dof$. Moreover, the decreasing margin of $\\chi _{min}^{2}$ and $\\chi _{min}^{2}/dof$ for these 10 DE models are almost same.  To further verify this conclusion, we also adopt the method of random truncation used in \\cite{ness07,wei10}. The method is very simple. First, we random select the outliers from the full Union dataset and the full Constitution dataset, respectively. Notice that the numbers of data points in these random ``UnionOut'' subsets are same as that in the original UnionOut samples, while the numbers of data points in these random ``ConstitutionOut'' subsets are same as that in the original ConstitutionOut samples. As in \\cite{ness07,wei10}, 500 random UnionOut subsets and 500 random ConstitutionOut subsets are selected. By subtracting these outliers from the Union dataset and the Constitution dataset, respectively, 500 random UnionT subsets and 500 random ConstitutionT subsets are obtained. Then, we perform best-fit analysis on the HDE model by using these random truncated SNIa subsets. Using the SNIa data only, we get the Mean of $\\chi_{min}^{2}$ for 500 random UnionOut subsets and 500 random ConstitutionOut subsets, respectively. Notice that the full Union set gives a $\\chi_{min}^{2}=310.827$ and the full Constitution set gives a $\\chi_{min}^{2}=465.769$, the differences between the $\\chi_{min}^{2}$ of the full SNIa dataset and the Mean of $\\chi_{min}^{2}$ of these random truncated SNIa subsets are also obtained. Next, we make a comparison for the original truncated SNIa subset and the random truncated SNIa subsets. The results are shown in table \\ref{table5}. It is found that the $\\chi_{min}^{2}$ of the HDE model can reduce 5 by dropping 1 data point in the original truncated SNIa subset, but can only reduce 1 by dropping 1 data point in the random truncated SNIa subset. Therefore, the original UnionOut and ConstitutionOut subsets are systematically different from the full Union and Constitution datasets, and the original truncated supernova datasets provide a significantly better model fitting than the full SNIa datasets.   \\begin{table} \\caption{A comparison for the original truncated SNIa subset and the random truncated SNIa subsets. Notice that the full Union set gives a $\\chi_{min}^{2}=310.827$ and the full Constitution set gives a $\\chi_{min}^{2}=465.769$. Here $\\Delta N$ is the number of data points in outliers, $\\overline{{\\chi^2}}$ is the Mean of $\\chi_{min}^{2}$ for those random truncated SNIa subsets, and $\\Delta \\chi^2$ is the differences between the $\\chi_{min}^{2}$ of the full SNIa dataset and the $\\chi_{min}^{2}$ (or the Mean) of the random truncated SNIa subsets.} \\begin{center} \\label{table5} \\begin{tabular}{ccccc} \\hline\\hline ~~~The subset of outliers~~~ & ~~~$\\Delta N$~~~ & ~~~$\\chi_{min}^{2}$ or $\\overline{{\\chi^2}}$~~~ & ~~~$\\Delta \\chi^2$~~~ & ~~~$\\Delta \\chi^2/\\Delta N$~~~ \\\\ \\hline\\hline ~~~The original UnionOut subset~~~ & ~~~21~~~ & ~~~$204.070$~~~ & ~~~$106.757$~~~ & ~~~$5.084$~~~ \\\\ \\hline ~~~500 random UnionOut subsets~~~ & ~~~21 (Mean)~~~ & ~~~$289.617$ (Mean)~~~ & ~~~$21.210$ (Mean)~~~ & ~~~$1.010$ (Mean)~~~ \\\\ \\hline ~~~The original ConstitutionOut subset~~~ & ~~~34~~~ & ~~~$269.113$~~~ & ~~~$199.656$~~~ & ~~~$5.784$~~~ \\\\ \\hline ~~~500 random ConstitutionOut subsets~~~ & ~~~34 (Mean)~~~ & ~~~$425.763$ (Mean)~~~ & ~~~$40.006$ (Mean)~~~ & ~~~$1.177$ (Mean)~~~ \\\\ \\hline\\hline \\end{tabular} \\end{center} \\end{table}   In \\cite{wei10}, by performing the best-fit analysis on the CPL model, Wei argued that the UnionT and the ConstitutionT datasets are fully consistent with the other cosmological observations. To verify this conclusion, more theoretical models should be taken into account. Here we study the HDE model as an example. In Fig.\\ref{fig1}, we plot the $1\\sigma$ and the $2\\sigma$ confidence level (CL) contours for the HDE model, where the Union and the UnionT datasets are used, respectively. From this figure, we find that the best-fit point for the Union data is outside the $2\\sigma$ confidence region given by the combined Union+BAO+CMB data, while the best-fit point for the UnionT data is inside the $1\\sigma$ confidence region given by the combined UnionT+BAO+CMB data. This means that the UnionT dataset is very useful to remove the tension between SNIa data and other cosmological observations. In addition, we also use the Constitution and the ConstitutionT datasets to plot the CL contours for the HDE model in Fig.\\ref{fig2}. It is found that the best-fit point for the Constitution data is outside the $2\\sigma$ confidence region given by the combined Constitution+BAO+CMB data, while the best-fit point for the ConstitutionT data is very close to the best-fit point for the combined ConstitutionT+BAO+CMB data. This means that the ConstitutionT dataset is also very helpful to remove the tension. Therefore, we conclude that the truncated SNIa datasets can remove the tension between SNIa data and other cosmological observations.   \\begin{figure} \\includegraphics[scale=0.7, angle=0]{HDE_Union.eps} \\includegraphics[scale=0.7, angle=0]{HDE_Union_T.eps} \\caption{\\label{fig1} The $1\\sigma$ and the $2\\sigma$ CL contours for the HDE model. The left panel is plotted by using the Union dataset, while the right panel is plotted by using the UnionT dataset. For both these two panels, the blue dashed lines correspond to the constraints given by the SNIa data only, and the red solid lines correspond to the constraints given by the combined SNIa+BAO+CMB data. Moreover, we also plot the best-fit point for the SNIa data (blue point) and the best-fit point for the combined SNIa+BAO+CMB data (red star). It should be pointed out that, in the left panel, the best-fit point for the SNIa data is outside the $2\\sigma$ confidence region given by the combined SNIa+BAO+CMB data, while in the right panel, the best-fit point for the SNIa data is inside the $1\\sigma$ confidence region given by the combined SNIa+BAO+CMB data. This means that the UnionT dataset is very useful to remove the tension between SNIa data and other cosmological observations.} \\end{figure}    \\begin{figure} \\includegraphics[scale=0.7, angle=0]{HDE_CFA.eps} \\includegraphics[scale=0.7, angle=0]{HDE_CFA_T.eps} \\caption{\\label{fig2}  The same as in Fig.\\ref{fig1}, except for the cases of the Constitution and the ConstitutionT datasets. It should be mentioned that, in the left panel, the best-fit point for the SNIa only is outside the $2\\sigma$ confidence region given by the combined SNIa+BAO+CMB data, while in the right panel, the best-fit point for the SNIa only is very close to the best-fit point for the combined SNIa+BAO+CMB data. This means that the ConstitutionT dataset is very helpful to remove the tension between SNIa data and other cosmological observations.} \\end{figure}       (2) The CMB data is very helpful to break the degeneracy among different parameters, and plays a very important role in distinguishing different DE models.   As an example, we will compare the HDE model with the RDE model. By using the Constitution SNIa data alone, the CMB data alone, and the combined Constitution+CMB+BAO data, respectively, we plot the $1\\sigma$ and the $2\\sigma$ CL contours for the HDE and the RDE model in Fig.\\ref{fig3}. From this figure, we find that the shapes of CL contours given by the CMB data alone are quite different from that given by the SNIa data alone, and the CMB data is very helpful to break the degeneracy among different parameters. As shown in the left panel, the $1\\sigma$ CL contour of the HDE model given by the CMB data intersects to the $1\\sigma$ CL contour of the HDE model given by the SNIa data; while in the right panel, the $1\\sigma$ CL contour of the RDE model given by the CMB data does not intersect to the $1\\sigma$ CL contour of the RDE model given by the SNIa data. This fact explains why the HDE model performs much better in fitting the combined Constitution+BAO+CMB data than the RDE model. Besides, we also check the effect of the BAO data, and find that the $1\\sigma$ confidence region of the HDE model given by the BAO data alone can completely cover that given by the SNIa data. So the constraint given by the BAO data is quite weaker than that given by the SNIa and the BAO data. This conclusion can be further verified by using table \\ref{table3}. As seen in table \\ref{table3}, without adopting the CMB data, it is very difficult to distinguish the HDE Model from the RDE model. After adding the CMB data, it is seen that the $\\chi_{min}^{2}$ of the HDE Model given by the SNIa+CMB data is 6.811 smaller than the $\\chi_{min}^{2}$ of the RDE Model given by the SNIa+CMB data, while the $\\chi_{min}^{2}$ of the HDE Model given by the combined SNIa+CMB+BAO data is 7.082 smaller than the $\\chi_{min}^{2}$ of the RDE Model given by the combined SNIa+CMB+BAO data.    \\begin{figure} \\includegraphics[scale=0.7, angle=0]{HDE_All_CFA.eps} \\includegraphics[scale=0.7, angle=0]{RDE_All_CFA.eps} \\caption{\\label{fig3} The $1\\sigma$ and the $2\\sigma$ CL contours for the HDE and the RDE model. The Constitution dataset is used to plot this figure. The left panel is plotted by using the HDE model, while the right panel is plotted by using the RDE model. For both these two panels, The blue dashed lines correspond to the constraints given by the SNIa data only, the black dotted lines correspond to the constraints given by the CMB data only, the red solid lines correspond to the constraints given by the combined SNIa+BAO+CMB data, and the red stars denote the best-fit point for the combined SNIa+BAO+CMB data. Since the shapes of CL contours given by the CMB data alone are quite different from that given by the SNIa data alone, the CMB data is very helpful to break the degeneracy among different parameters. As shown in the left panel, the $1\\sigma$ CL contour of the HDE model given by the CMB data intersects to the $1\\sigma$ CL contour of the HDE model given by the SNIa data; while in the right panel, the $1\\sigma$ CL contour of the RDE model given by the CMB data does not intersect to the $1\\sigma$ CL contour of the RDE model given by the SNIa data. This fact explains why the HDE model performs much better in fitting the combined Constitution+BAO+CMB data than the RDE model. Besides, we also check the effect of the BAO data, and find that the $1\\sigma$ confidence region of the HDE model given by the BAO data alone can completely cover that given by the SNIa data. For simplicity, the CL contours given by the BAO data are not plotted in this figure.} \\end{figure}  In addition, by using the ConstitutionT SNIa data alone, the CMB data alone, and the combined ConstitutionT+CMB+BAO data, respectively, we also plot the $1\\sigma$ and the $2\\sigma$ CL contours for the HDE and the RDE model in Fig.\\ref{fig4}. Again, we see that the $1\\sigma$ CL contour of the HDE model given by the CMB data intersects to the $1\\sigma$ CL contour of the HDE model given by the SNIa data, while the $1\\sigma$ CL contour of the RDE model given by the CMB data does not intersect to the $1\\sigma$ CL contour of the RDE model given by the SNIa data. This fact explains why the HDE model performs much better in fitting the combined ConstitutionT+BAO+CMB data than the RDE model. Therefore, the CMB data is very helpful to break the degeneracy among different parameters, and plays a very important role in distinguishing different DE models.    \\begin{figure} \\includegraphics[scale=0.7, angle=0]{HDE_All_CFA_T.eps} \\includegraphics[scale=0.7, angle=0]{RDE_All_CFA_T.eps} \\caption{\\label{fig4} The same as in Fig.\\ref{fig3}, except for the case of the ConstitutionT dataset. Notice that the shapes of CL contours given by the CMB data alone are quite different from that given by the SNIa data alone. As shown in the left panel, the $1\\sigma$ CL contour of the HDE model given by the CMB data intersects to the $1\\sigma$ CL contour of the HDE model given by the SNIa data; while in the right panel, the $1\\sigma$ CL contour of the RDE model given by the CMB data does not intersect to the $1\\sigma$ CL contour of the RDE model given by the SNIa data. This fact explains why the HDE model performs much better in fitting the combined ConstitutionT+BAO+CMB data than the RDE model.} \\end{figure}   (3) The current observational data are still too limited to distinguish all DE models.  First, we discuss three kinds of two-parameter models: the XCDM model, the CG model, and the HDE model. As seen in table \\ref{table3}, the $\\chi_{min}^{2}$ of these 3 models given by the combined Constitution+BAO+CMB data are 466.947, 466.891, and 467.384, respectively. That is to say, after taking into account the CMB data, the differences of the $\\chi_{min}^{2}$ of these 3 models are still smaller than 1. Notice that this result also holds true in table \\ref{table1}, table \\ref{table2}, and  table \\ref{table4}. Therefore, one cannot judge which DE model is better.  Then, we discuss  three kinds of three-parameter models: the LP model, the CPL model, and the GCG model. As seen in table \\ref{table3}, the $\\chi_{min}^{2}$ of these 3 models given by the combined Constitution+BAO+CMB data are 466.910, 466.902, 466.895, respectively, and the differences of the $\\chi_{min}^{2}$ of these 3 models are even smaller than 0.1. Since this result also holds true in table \\ref{table1}, table \\ref{table2}, and  table \\ref{table4}, it is also very difficult to distinguish these 3 models. Therefore, to distinguish DE models better, more high-quality observational data are needed."
    },
    "0910/0910.2712_arXiv.txt": {
        "abstract": "We combine a cosmological reionization simulation with box size of 100$\\hMpc$ on a side and a Monte Carlo \\lya radiative transfer code to model Lyman Alpha Emitters (LAEs) at $z\\sim$5.7. The model introduces \\lya radiative transfer as the single factor for transforming the intrinsic \\lya emission properties into the observed ones. Spatial diffusion of \\lya photons from radiative transfer results in extended \\lya emission and only the central part with high surface brightness can be observed. Because of radiative transfer, the appearance of LAEs depends on density and velocity structures in circumgalactic and intergalactic media as well as the viewing angle, which leads to a broad distribution of apparent (observed) \\lya luminosity for a given intrinsic \\lya luminosity. Radiative transfer also causes frequency diffusion of \\lya photons. The resultant \\lya line is asymmetric with a red tail. The peak of the \\lya line shifts towards longer wavelength and the shift is anti-correlated with the apparent to intrinsic \\lya luminosity ratio. The simple radiative transfer model provides a new framework for studying LAEs. It is able to explain an array of observed properties of $z\\sim$5.7 LAEs in Ouchi et al. (2008), producing \\lya spectra, morphology, and apparent \\lya luminosity function (LF) similar to those seen in observation. The broad distribution of apparent \\lya luminosity at fixed UV luminosity provides a natural explanation for the observed UV LF, especially the turnover towards the low luminosity end. The model also reproduces the observed distribution of \\lya equivalent width (EW) and explains the deficit of UV bright, high EW sources. Because of the broad distribution of the apparent to intrinsic \\lya luminosity ratio, the model predicts effective duty cycles and \\lya escape fractions for LAEs. ",
        "introduction": "More than four decades ago, \\citet{Partridge67} proposed that prominent \\lya emission reprocessed from ionizing photons of young stars in galaxies can be used to detect high-redshift galaxies. The first successful detections of high-redshift \\lya emitting galaxies, or \\lya emitters (LAEs), were made $\\sim$ 30 years later \\citep[e.g.,][]{Hu96,Cowie98,Dey98,Hu98,Hu99}. Recently, important advances have been made on the observational front to detect LAEs at $z\\gtrsim 6$ \\citep[e.g.,][]{Hu98,Hu02,Hu04,Hu05,Hu06,Rhoads03,Malhotra04,Horton04,Stern05, Kashikawa06,Shimasaku06,Iye06,Cuby07,Ouchi07,Ouchi08,Stark07a,Nilsson07, Willis08,Ota08}.  LAEs can be efficiently detected through narrow-band imaging or with integral-field-units (IFU) spectroscopy. Owing to the high efficiency of target detection, LAEs naturally become objects for large surveys of high-redshift galaxies. Besides providing clues to the formation and evolution of galaxies at the time when the universe was still young, LAEs are an important tracer of the large-scale structure. The clustering of LAEs may be used to constrain cosmological parameters. In particular, the large-volume surveys such as the Hobby-Eberly Telescope Dark Energy Experiment (HETDEX; \\citealt{Hill08}) will enable the detection of the baryon acoustic oscillations (BAO) features \\citep[e.g.,][]{Eisenstein05} in the LAE power spectrum. The BAO and the shape of the power spectrum can be used to measure the expansion history of the universe at early epochs ($z\\sim 3$), which constrains the evolution of dark energy and the curvature of the universe.  LAEs are also a key probe of the high-redshift intergalactic medium (IGM), especially across the reionization epoch. The use of LAEs to learn about reionization has been the subject of intense study \\citep[e.g.,][]{Rees98,Miralda98,Haiman99,Santos04,Haiman05,Dijkstra07a, Wyithe07}. Suitably devised statistics, including luminosity function (LF) and correlation functions of LAEs, can be used to constrain the neutral fraction of the IGM during reionization \\citep[]{Malhotra04,Haiman05,Kashikawa06,Furlanetto06,Dijkstra07b,McQuinn07, Mesinger08,Iliev08,Dayal08,Dayal09}. By comparing the LFs of $z\\sim 5.7$ and $z\\sim 6.5$ LAEs, \\citet{Malhotra04} conclude that reionization was largely complete at $z\\sim 6.5$ (also see \\citealt{Dijkstra07b}). \\citet{McQuinn07} show that, with the angular correlation function of the 58 available $z\\sim 6.6$ LAEs in the Subaru Deep Field \\citep{Kashikawa06}, limits may be placed on the IGM neutral fraction, favoring a fully ionized universe at $z\\sim 6.6$.   However, none of the previous work of LAEs mentioned above used reionization simulations with concurrent treatment of hydrodynamics plus radiative transfer of ionizing photons and \\lya photons. Hydrodynamic and radiative transfer simulations provide realistic neutral gas distributions, and \\lya radiative transfer yields detailed properties of the \\lya emission. Realistic \\lya radiative transfer calculations have been applied to high-redshift LAEs in cosmological simulations \\citep[e.g.,][]{Tasitsiomi06}. The application, however, is limited to a few individual sources, which do not form a sample for statistical study.  \\citet{McQuinn07} and \\citet{Iliev08} studied a sample of LAEs in reionization simulations with cosmological volume. However, the radiative transfer of \\lya photons is treated in a simplistic way in their study: the observed \\lya spectrum is modeled as the intrinsic line profile modified by $\\exp(-\\tau_\\nu)$, where $\\tau_\\nu$ is the optical depth at frequency $\\nu$ along the line of sight. Although this $\\exp(-\\tau_\\nu)$ model can yield insights into the properties of the observed \\lya emission, such as the effect of IGM on the observability of LAEs, it is far from a complete description of the radiative transfer of \\lya photons. First, during the propagation, \\lya photons experience frequency diffusion, which is neglected by the simple $\\exp(-\\tau_\\nu)$ model. The $\\exp(-\\tau_\\nu)$ model removes \\lya photons at a given frequency according to the \\lya optical depth, and no frequency change occurs for any \\lya photon, therefore it does not yield correct \\lya spectra. Second, the simple $\\exp(-\\tau_\\nu)$ model does not account for the spatial diffusion of \\lya photons either. LAEs in this model appear as point sources in \\lya and there is no surface brightness information. Even if \\lya photons start from a point source, spatial diffusion due to radiative transfer would lead to an extended source. Observationally, LAEs indeed appear to be extended and they are defined by a surface brightness threshold in the narrow-band image \\citep[e.g.,][]{Ouchi08}. Therefore, although the simple $\\exp(-\\tau_\\nu)$ model may provide useful insight, it likely falls short for predicting the detailed properties of the observed \\lya emission from LAEs.  To correctly understand high-redshift LAEs and use them for cosmological study, a full calculation of radiative transfer of \\lya photons for a large sample of LAEs in cosmological reionization simulation is necessary, as will be evident later. In this work, we aim to perform detailed radiative transfer calculation of \\lya photons \\citep{Zheng02} from LAEs in a self-consistent fashion through radiation-hydrodynamic reionization simulations \\citep{Trac08}. For this paper, we focus on studying statistical properties of $z\\sim 5.7$ LAEs and show how the radiative transfer calculation aids our understanding of the observed properties of LAEs. The clustering properties of LAEs from this study will be presented in another paper (Paper II; Zheng et al. in prep.). The paper is organized as follows. In \\S~2 we review the cosmological reionization simulation used in this work and in \\S~3 we describe the \\lya radiative transfer calculation. In \\S~4, we study in details the \\lya emission from an individual source chosen from the simulation box to gain a general view of the effect of \\lya radiative transfer on the appearance of LAEs. Then, we present the statistical properties of LAEs in \\S~5, including their spectra and luminosity, from our modeling of an ensemble of sources in the simulation box. We compare our modeling results with observations for $z\\sim 5.7$ LAEs and discuss the implications in our understanding of LAEs. \\S~7 is devoted to identifying important physical factors in shaping the observed \\lya emission of LAEs. We summarize and discuss the results in \\S~8. ",
        "conclusions": "\\subsection{Summary of Main Results}  We perform a full \\lya radiative transfer calculation with a Monte Carlo code \\citep{Zheng02} to study LAEs in a cosmological volume. The LAE sources and the physical properties of neutral hydrogen gas are taken from the $z\\sim 5.7$ outputs of a cosmological reionization simulation \\citep{Trac08}, which solves the coupled evolution of the dark matter, baryons, and ionizing radiation in a box of 100$\\hMpc$ (comoving) on a side. The large volume of the simulation allows a statistical study of $z\\sim 5.7$ LAEs. Radiative transfer of \\lya photons in the IGM environment around LAEs, which leads to both frequency and spatial diffusion of \\lya photons, turns out to play a crucial role in determining the observability of LAEs and in understanding the observed properties of LAEs.  Although the radiative transfer calculation is computationally costly, the LAE model we present in this paper is rather simple. The UV or intrinsic \\lya luminosity is assumed to be proportional to the SFR, which is tightly coupled to halo mass in the reionization simulation we use. That is, we essentially adopt a constant mass to light ratio, where mass is halo mass and light is either UV or \\lya. All we do is to add the physics of \\lya radiative transfer into the model to obtain the observed properties of LAEs. That is, we introduce the radiative transfer of \\lya photons in the IGM as the single factor responsible for transforming the intrinsic \\lya emission properties to the observed ones. Our model produces IFU-like data cube that covers the extent of the simulation box, which allows mock observations to be made. With the \\lya image contracted from this data cube, we follow typical observational procedures \\citep[e.g.,][]{Ouchi08} to identify LAEs and then extract their \\lya spectra.  Initially \\lya photons are produced inside the star formation region. Therefore the intrinsic \\lya sources are expected to be similar in size as the UV sources, which are compact ($\\lesssim 1$ kpc; \\citealt{Taniguchi09}). We find that an intrinsically point-like \\lya source becomes extended as a consequence of resonant scatterings of \\lya photons (spatial diffusion). The scatterings of \\lya photons do not destroy them and all \\lya photons escape in the end. However, observationally, only the central part of the extended source can be detected as a consequence of the limit set by the surface brightness threshold. The scatterings of \\lya photons also cause the frequency of \\lya photons to change (frequency diffusion). The resultant \\lya spectra from the central aperture do not have a simple relation to the initial profile, which is assumed to be Gaussian in our model. Our results from full radiative transfer calculations show a clear difference from a simple treatment of \\lya radiative transfer, namely the $\\exp(-\\tau_\\nu)$ model, widely adopted in previous work, which modifies the intrinsic \\lya spectrum by multiplying the line-of-sight transmission determined by the optical depth at each frequency.  The observed \\lya spectrum of an LAE in our model shows a clear asymmetry, skewed towards red. Although the $\\exp(-\\tau_\\nu)$ model produces the same qualitative feature, the predicted line profile, the frequency shift, and the total flux are all significantly different from our results. While the spectrum of the $\\exp(-\\tau_\\nu)$ model we present is essentially the intrinsic one truncated below a certain wavelength (but see Figures 14 and 15 in \\citealt{Iliev08} for more complex line shapes, probably caused by different assumptions in the $\\exp(-\\tau_\\nu)$ model), the observed \\lya spectrum in our model can have contributions from photons with frequency much redder than initial photons, a result of the scattering-caused frequency diffusion. We find that the redward shift of the Lya line induced by radiative transfer is usually a few times the intrinsic line width, with a distribution that peaks at about three times. The asymmetry and shift of the \\lya line do not indicate the presence of any winds, but they arise from the structure of the halo infall and Hubble expansion around the sources. If one were to infer the wind velocity, if there is any, from comparing the relative shift in the \\lya line and an optically-thin line, one has to keep in mind the \\lya radiative transfer effect. For example, the observationally inferred velocity of the receding winds would be overestimated by $\\sim 100\\kms$ or more if the effect is not taken into account.  As a consequence of the frequency diffusion and spatial diffusion, our model predicts a much higher observed \\lya flux than the $\\exp(-\\tau_\\nu)$ model. At a fixed intrinsic \\lya luminosity (i.e., fixed host halo mass), the observed (apparent) luminosity is broadly distributed. The shift in the \\lya line peak and the ratio of the apparent to intrinsic \\lya luminosity appear to be anti-correlated. The distributions of the line peak shift and the ratio of apparent to intrinsic \\lya luminosity, and their correlation, all result from the dependence of the \\lya radiative transfer on the IGM environments around sources. It would be interesting and extremely useful if we could make use of the full information in the observed \\lya properties to infer the intrinsic ones, and we reserve such an investigation for future work. Although our model predicts a much higher observed \\lya flux than the $\\exp(-\\tau_\\nu)$ model, it still leads to a highly suppressed \\lya flux, compared with the intrinsic one. The suppression factor depends on the assumed line width of the intrinsic \\lya spectra. For the line width assumed in our model (given by halo virial temperature), we find that, with respect to the intrinsic \\lya LF of LAEs, the observed (apparent) \\lya LF shift towards the low luminosity end by roughly one order of magnitude in luminosity. For comparison, the $\\exp(-\\tau_\\nu)$ model would shift by two orders of magnitude in luminosity.  We make comparisons between the $z\\sim 5.7$ LAEs in our model and those observed in SXDS \\citep{Ouchi08}. The sizes, morphologies, \\lya line profiles of the model LAEs are remarkably similar to the observed ones. For the \\lya LF, UV LF, and \\lya EW distribution, our model can successfully reproduce the observations and provide physical explanations for various observed features.  After an overall adjustment of a factor of $\\sim 5$ in luminosity, the \\lya LF of model LAEs matches well with observation. The adjustment reflects our incomplete knowledge in the stellar IMF at high redshift, the uncertainty in the model SFR, and the lack of information on the intrinsic \\lya line profile. According to our model, there is no one-to-one map between the intrinsic and the observed \\lya luminosity. In other words, there is a large scatter in the relation between the apparent and intrinsic luminosities. At a fixed observed luminosity, LAEs can differ by one order of magnitude in the intrinsic luminosity (Fig.~\\ref{fig:LL}$c$). This large scatter has to be taken into account when interpreting the observed \\lya LF and linking the observed LAEs to their host halos.  For the UV LF of observed LAEs, our model prediction shows a good agreement with observation. In particular, the turnover of the UV LF towards the low luminosity end seen in high-$z$ ($z\\sim$ 3--6) LAEs are well reproduced. The key to interpret the shape of the UV LF is that observed LAEs are sources with observed (apparent) \\lya luminosity above certain threshold. The turnover reflects that for LAEs with low UV luminosity (or low intrinsic \\lya luminosity, or low halo mass), the probability for the observed \\lya luminosity to exceed the observation threshold is low, a consequence of the broad distribution of apparent \\lya luminosity at a given intrinsic \\lya luminosity. The full UV LF for sources in our model (i.e., without imposing the observation \\lya luminosity threshold) agrees well with the nonduplicated sum of the observed UV LFs of LAEs and $i$--dropout galaxies at $z\\sim 6$.  The observed distribution of \\lya EW as a function of UV luminosity is also reproduced in the model. We note that in our model all the sources have the same {\\it intrinsic} EW and the distribution of the {\\it observed} values of EW at fixed UV luminosity is purely caused by the environment-dependent radiative transfer effect. At a fixed UV luminosity, the distribution of observed (apparent) EW is a decreasing function toward high values. The observational trend of lacking UV bright, high EW sources \\citep[e.g.,][]{Ando06} is naturally explained by our model in that such sources lie in a low probability corner --- a combination of the drop of the UV LF toward high luminosity and the drop of the apparent EW distribution function toward high EW value. LAE surveys with large volume will test the interpretation.  Therefore, the observed properties of LAEs can be explained by simply invoking \\lya radiative transfer: the effects of the local IGM environment, depending mainly on the gas density and line-of-sight velocity and their line-of-sight gradients, lead to the distribution of observed \\lya emission properties at fixed intrinsic \\lya luminosity. This environmental selection also causes new features in the clustering of LAEs that we will study in Paper II.  \\subsection{Implications and Discussion}  Our interpretation of the observations of LAEs does not invoke any mass dependent dust absorption, which is in contrast to many previous models \\citep[e.g.,][]{Dayal10}. Uniformly distributed dust efficiently absorbs \\lya photons, since the large number of resonant scatterings increase the path length. There is much less attenuation when the dust is in gas clumps and \\lya photons bounce off the cloud surfaces \\citep{Neufeld91,Hansen06}. Optical, UV, and \\lya observations of local star-forming galaxies provide evidence that ISM kinematics and geometry play a more significant role than dust in affecting the \\lya emission \\citep[e.g.,][]{Giavalisco96,Keel05,Atek08, Atek09}. Our model successfully reproduces the observed UV LF of LAEs by incorporating only a mass independent effective UV extinction of at most 0.3 mag. We conclude that any mass dependent dust effects are not likely to play a substantial role to determine the observed properties of LAEs, compared to \\lya radiative transfer effects.  Our model also has important implications for the duty cycle and the \\lya escape fraction of LAEs. The theoretically predicted (intrinsic) \\lya LF, which essentially is the halo mass function, is substantially higher than the observed one. Two scenarios have been introduced to address this problem, the duty cycle and the \\lya escape fraction scenarios \\citep[e.g.,][]{Stark07b,Nagamine08}. In the duty cycle scenario, LAEs are short-lived and a fraction of all galaxies are active as LAEs at any given time, lowering the amplitude of the \\lya LF. In the \\lya escape fraction scenario, only a fixed fraction of \\lya photons escape from the source and the overprediction problem is solved by shifting the LF towards the low luminosity end. To conserve the number density of LAEs of a given sample, the masses of host halos in the duty cycle scenario would be on average lower than those in the escape fraction scenario. As a consequence, the clustering of LAEs would be different in the two scenarios, with a stronger clustering in the escape fraction scenario. \\citet{Nagamine08} find that LAE clustering measurements from observations are in favor of their duty cycle scenario.  In our model, \\lya photons all escape after a large number of scatterings. The \\lya escape fraction, in its literal meaning, is therefore unity. However, only the central part of the extended \\lya emission of LAEs can be observed, which gives rise to an {\\it apparent} or {\\it effective} \\lya escape fraction. Since the observed \\lya luminosity has a broad distribution at a fixed intrinsic \\lya luminosity, our model predicts a broad distribution of the effective \\lya escape fraction rather than a single value. In our model, no duty cycle parameter is introduced. Since halos of the same mass have similar SFR in our model, the corresponding intrinsic \\lya luminosities are the same, i.e., \\lya emission does not come from a fraction of halos. However, an {\\it apparent} or {\\it effective} duty cycle arises as a result of the selection effect caused by \\lya radiative transfer (a broad distribution of observed \\lya luminosity at a fixed intrinsic \\lya luminosity) and a \\lya luminosity threshold in observation. This can be seen from comparing the UV LF for all galaxies (dropout galaxies and LAEs) and that for LAEs (Fig.~\\ref{fig:uvLF}), which can be described as that at a fixed UV luminosity (or halo mass) only a fraction of all the galaxies are observed as LAEs. This effective duty cycle does not have the physical meaning in its original form. Moreover, it is not a constant, since it changes with UV luminosity (Fig.~\\ref{fig:uvLF}). \\citet{Tilvi09} present an LAE model in which \\lya luminosity (or SFR) is related to the halo mass accretion rate, rather than halo mass, and the model naturally gives rise to the duty cycle of LAEs. The duty cycle in their model, however, has its original meaning, in direct contrast with our model. Our model still ties the intrinsic \\lya luminosity (SFR) to halo mass and let the \\lya radiative transfer do the work of converting it to observed \\lya luminosity. Because of the large scatter between the observed \\lya luminosity and the intrinsic one (or halo mass) in our model, for a sample of LAEs above a \\lya luminosity threshold, some of them can reside in halos with mass smaller than the threshold mass above which halos have the same number density as LAEs. So the effective duty cycle in our model has the effect of lowering the clustering amplitude of LAEs.  The main uncertainties in our model are the stellar IMF, the SFR, and the intrinsic \\lya line profile. The first two are general uncertainties for any model. The IMF at high-$z$ is neither well constrained observationally nor well understood theoretically. The SFR in galaxy formation model is related to the complex gas physics that we do not have a satisfactory understanding. Changing IMF or SFR would change the details in the reionization process and therefore change the gas properties (e.g., neutral fraction and temperature distribution) at a fixed redshift. For the reionization history itself, the escape fraction of ionizing photons adds a further uncertainty. Even though we focus on LAEs at $z\\sim 5.7$, when reionization is almost complete, different reionization histories can still leave different imprints on the gas distribution. For example, the IGM temperature in a region is correlated to the time when this region is reionized and heated \\citep[e.g.,][]{Trac08}. A detailed study is needed to investigate the effect of inhomogeneous IGM temperature distribution on \\lya radiative transfer. To be fully self-consistent, for any change in the IMF, SFR, and escape fraction of ionizing photons, one has to re-run the reionization simulation to solve the density, velocity, and temperature distributions of neutral gas and then perform the \\lya radiative transfer calculation. If the IMF, SFR, and escape fraction of ionizing photons change in a way to maintain the same reionization history, the effect of the IMF and SFR change can be largely characterized by an overall scaling in UV or intrinsic \\lya luminosity and one does not need to redo the \\lya radiative transfer calculation. For simplicity and to avoid extensive computations for reionization and radiative transfer simulation, we adopt such a scenario in this paper. Changing the width of the intrinsic \\lya line profile leads to changes in the distribution of the apparent to intrinsic \\lya luminosity ratio. Although we cast the effect as an overall scaling in the apparent \\lya luminosity for the \\lya LF, the intrinsic \\lya line profile is important in many aspects of the \\lya observation (image, spectra, etc) and its effect deserves a detailed investigation. \\lya radiative transfer calculation with high resolution hydrodynamic simulations for individual LAEs are necessary to shed light on the intrinsic \\lya line profile \\citep[e.g.,][]{Laursen07}. Changing the IMF affects the flux of ionizing photons more than that of the $\\sim$ 1500\\AA UV photons, which leads to a change in the ratio of the intrinsic \\lya luminosity to UV luminosity, or the intrinsic \\lya EW. Intrinsic \\lya line profile and dust play a role in converting the intrinsic EW to apparent EW, with the former affecting the \\lya luminosity and the latter adding extinction to UV luminosity. Therefore, the full distribution of observed \\lya EW and UV luminosity of LAEs can give constraints on the IMF, SFR, dust, and intrinsic \\lya profile.  Although we have identified key environment factors in shaping the observational properties of LAEs, the dependence of radiative transfer on environments deserves a further study to understand the details of \\lya scatterings in the surrounding regions of LAEs.  Our \\lya radiative transfer calculation relies on the gas distribution and properties from the cosmological reionization simulation. The radiative transfer of ionizing photons with a ray tracing algorithm is crucial in determining the state of gas. We have tested our LAE model for a reionization simulation with improved ray tracing algorithm (Trac et al. in prep.) in a small box (25$\\hMpc$ on a side). We find that the results presented in this paper are robust.  \\lya radiative transfer through the surrounding circumgalactic and intergalactic media is a physical process that likely plays an important role in galaxies at all redshifts. It has to be taken into account for modeling LAEs and for interpreting observations. Our model is rather {\\it simple} and can naturally explain an array of observations of LAEs, which make it extremely attractive. It is interesting to see how well it does in interpreting observations of LAEs at lower redshifts (e.g., $z\\sim 3$). We also plan to apply it to the era of the late stage of reionization to study how to use LAEs to constrain reionization."
    },
    "0910/0910.2238_arXiv.txt": {
        "abstract": "We investigate how the removal of interstellar material by stellar feedback limits the efficiency of star formation in molecular clouds and how this determines the shape of the mass function of young star clusters. In particular, we derive relations between the power-law exponents of the mass functions of the clouds and clusters in the limiting regimes in which the feedback is energy-driven and momentum-driven, corresponding to minimum and maximum radiative losses, and likely to bracket all realistic cases. We find good agreement between the predicted and observed exponents, especially for momentum-driven feedback, provided the protoclusters have roughly constant mean surface density, as indicated by observations of the star-forming clumps within molecular clouds. We also consider a variety of specific feedback mechanisms, concluding that \\hii\\ regions inflated by radiation pressure predominate in massive protoclusters, a momentum-limited process when photons can escape after only a few interactions with dust grains. We show in this case that the star formation efficiency depends on the masses and sizes of the protoclusters only through their mean surface density, thus ensuring consistency between the observed exponents of the mass functions of the clouds and clusters. Our numerical estimate of this efficiency is also consistent with observations. ",
        "introduction": "Most stars form in protoclusters in dense molecular clumps (\\citealt{lada03}; \\citealt{mckee07b}). The energy and momentum injected by young stars then removes the remaining interstellar material (ISM), thus ending further star formation and reducing the gravitational binding energy of the protoclusters. This feedback limits the efficiency of star formation---the ratio of final stellar mass to initial interstellar mass---to only $20-30\\%$, and leaves many protoclusters unbound, with their constituent stars free to disperse. Even those protoclusters that survive will lose some stars by ISM removal and subsequent processes.  Two of the best probes of these formation and disruption processes are the mass functions of molecular clouds and young star clusters, defined as the number of objects per unit mass, $\\psi(M) \\equiv dN/dM$. For molecular clouds, the best-studied galaxies are the Milky Way and the Large Magellanic Cloud (LMC), while for star clusters, they are the Antennae and the LMC. In these and other cases, the observed mass functions can be represented by power laws, $\\psi(M) \\propto M^{\\beta}$, from $10^4M_{\\odot}$ or below to $10^6M_{\\odot}$ or above. Giant molecular clouds (GMCs) identified in CO surveys have $\\beta \\approx -1.7$ \\citep{rosolowsky05b, blitz07a, fukui08a}. This exponent is also found for massive self-gravitating clumps within GMCs, the formation sites of star clusters, whether they are identified by CO emission \\citep{bertoldi92} or higher-density tracers such as C$^{18}$O, $^{13}$CO, and thermal dust emission \\citep{reid06b, munoz07a, wong08a}. Young star clusters have $\\beta \\approx -2.0$ \\citep{elmegreen97a, mckee97, zhang99b, dowell08a, fall09a, chandar09a}. The similar exponents for clouds and clusters indicate that the efficiency of star formation and probability of disruption are at most weak functions of mass. This conclusion is reinforced by the fact that $\\beta$ is the same for $10^7-10^8$ yr-old clusters as it is for $10^6-10^7$ yr-old clusters \\citep{zhang99b, fall09a, chandar09a}.  These empirical results may at first seem puzzling. Low-mass protoclusters have lower binding energy per unit mass and should therefore be easier to disrupt than high-mass protoclusters. Indeed, several authors have proposed that feedback would cause a bend in the mass function of young clusters at $M \\sim 10^5M_{\\odot}$, motivated in part by the well-known turnover in the mass function of old globular clusters \\citep{kroupa02a, baumgardt08a, parmentier08a}. For young clusters, such a feature is not observed (as noted above), while for globular clusters, it arises from almost any initial conditions as a consequence of stellar escape driven by two-body relaxation over $\\sim 10^{10}$~yr \\citep[and references therein]{fall01a, mclaughlin08a}. Nevertheless, we are left with an important question: What are the physical reasons for the observed similarity of the mass functions of molecular clouds and young star clusters?  The goal of this Letter is to answer this question. In Section~\\ref{energymomentum}, we derive some general relations between the mass functions of clouds and clusters. In Section~\\ref{sec:efficiency}, we review a variety of specific feedback processes and estimate the star formation efficiency for radiation pressure, the dominant process in massive, compact protoclusters. We summarize in Section~\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  This Letter contains two main results. The first is the relation between the power-law exponents of the mass functions of molecular clouds and young star clusters, $\\beta_o$ and $\\beta_*$, in the limiting regimes in which stellar feedback is energy-driven and momentum-driven, Equations (\\ref{betastar:energy}) and (\\ref{betastar:momentum}), which bracket all realistic cases. The predicted $\\beta_*$ depends significantly on the initial size-mass relation of the protoclusters. We find good agreement between the predicted and observed $\\beta_*$, especially for momentum-driven feedback, for $\\Sigma \\propto M/R_h^2 \\approx {\\rm constant}$, the relation indicated by observations of gas-dominated protoclusters. In this case, the star formation efficiency is independent of protocluster mass, ensuring that the fraction of clusters that remain gravitationally bound following ISM removal is also independent of mass.  The second main result is an estimate of the star formation efficiency in protoclusters regulated by radiation pressure, Equations (\\ref{efficiency}) and (\\ref{Sigma_crit}). This is likely to be the dominant feedback process in massive protoclusters. We show that ${\\cal E}$ depends on $M$ and $R_h$ only through the mean surface density $\\Sigma$, which in turn guarantees consistency between the observed power-law exponents of the mass functions of molecular clouds and young star clusters according to our general relations. For $\\Sigma \\sim 1$~g~cm$^{-2}$, we estimate ${\\cal E} \\sim 0.3$, in satisfactory agreement with observations."
    },
    "0910/0910.5418_arXiv.txt": {
        "abstract": "{This paper reports the results obtained on the photometric redshifts measurement and accuracy, and cluster tomography in the ESO Distant Cluster Survey (EDisCS) fields.} {We present the methods used to determine photometric redshifts to discriminate between member and non-member galaxies and reduce the contamination by faint stars in subsequent spectroscopic studies.} {Photometric redshifts were computed using two independent codes both based on standard spectral energy distribution (SED) fitting methods ($Hyperz$ and G. Rudnick's code).  Simulations were used to determine the redshift regions for which a reliable determination of photometric redshifts was expected. The accuracy of the photometric redshifts was assessed by comparing our estimates with the spectroscopic redshifts of $\\sim 1400$ galaxies in the $0.3\\le z\\le 1.0$ domain. The accuracy expected for galaxies fainter than the spectroscopic control sample was estimated using a degraded version of the photometric catalog for the spectroscopic sample.  } {The accuracy of photometric redshifts is typically $\\sigma(\\Delta z/(1+z))\\sim 0.05 \\pm 0.01$, depending on the field, the filter set, and the spectral type of the galaxies. The quality of the photometric redshifts degrades by a factor of two in $\\sigma(\\Delta z/(1+z))$ between the brightest ($I\\ltapprox$22) and the faintest ($I\\sim$24-24.5) galaxies in the EDisCS sample. The photometric determination of cluster redshifts in the EDisCS fields using a simple algorithm based on \\zphot is in excellent agreement with the spectroscopic values, such that $\\delta z \\sim$0.03-0.04 in the high-z sample and $\\delta z \\sim$0.05 in the low-z sample, i.e. the \\zphot cluster redshifts are at least a factor $\\sim(1+z)$ more accurate than the measurements of \\zphot for individual galaxies. We also developed a method that uses both photometric redshift codes jointly to reject interlopers at magnitudes fainter than the spectroscopic limit. When applied to the spectroscopic sample, this method rejects $\\sim 50-90\\%$ of all spectroscopically confirmed non-members, while retaining $\\gtrsim 90\\%$ of all confirmed members.} {Photometric redshifts are found to be particularly useful for the identification and study of clusters of galaxies in large surveys. They enable efficient and complete pre-selection of cluster members for spectroscopy, allow accurate determinations of the cluster redshifts based on photometry alone, and provide a means of determining cluster membership, especially for bright sources. } ",
        "introduction": " ",
        "conclusions": "Discussion and Conclusions}  We have used two independent codes to compute photometric redshifts: {\\it Hyperz} and GR code. In general, the two codes yield rather similar results, either on a cluster-by-cluster basis or as a function of the filter set and spectral type, of typically $\\sigma(\\Delta z/(1+z))\\sim 0.05$ to 0.06. {\\it Hyperz} results are found to be slightly more accurate than GR's ones in general, by $\\ltapprox$20\\% in $\\sigma(\\Delta z/(1+z))$. The quality achieved by both codes is consistent with the expectations derived from ``ideal'' simulations. An interesting trend is that the quality of both codes is highly correlated, in the sense that the highest and lowest quality results, in terms of $\\sigma(\\Delta z/(1+z))$, and systematics are found for the same clusters. This trend cannot be due to the use of an incomplete or imperfect template set, as suggested by other authors (Ilbert et al.\\ 2006), because in such a case, we should expect the same systematic behavior in all fields, given a filter set, as discussed in \\S~\\ref{simus}. In contrast, different systematics are observed in the different fields, which are found to be almost equal for the two independent \\zphot codes. This behavior suggests that the origin of the systematic errors is more likely to be associated with small residuals in the input photometry rather than the \\zphot templates and codes. Indeed, small zero-point shifts of $\\ltapprox$0.05 magnitudes cannot be excluded, in particular for the near-IR data.  Photometric redshifts are found to be particularly useful in the identification and study of galaxy clusters in large surveys. The determination of cluster redshifts in the EDisCS fields using a simple algorithm based on \\zphot is highly accurate. Indeed, the differences between photometric and spectroscopic values are found to be small, typically ranging between $\\delta z \\sim$0.03-0.04 in the high-z sample and $\\delta z \\sim$0.05 in the low-z sample. This is at least a factor $\\sim(1+z)$ more accurate than the determination of \\zphot for individual galaxies. The accuracy is more sensitive to the filter set used rather than the redshift of the cluster. The systematic lower quality results for the low-z sample was somewhat expected from the simulations presented in \\S~\\ref{simus}. Tomography based on \\zphot could be used in searches for clusters along the line-of-sight, using redshift steps optimized to be close in value to the typical difference between photometric and spectroscopic z$_{cluster}$ to maximize the contrast between members and non-member galaxies (in this case, $\\Delta z \\sim 0.05$).  The cluster membership criterion presented in Sect.~\\ref{member} has been used to extend the spectroscopic studies of cluster galaxies to fainter limits in magnitude (e.g. De Lucia et al.\\ 2004, White et al.\\ 2005, Clowe et al.\\ 2006, Poggianti et al.\\ 2006, De Lucia et al.\\ 2007, Desai et al.\\ 2007, Rudnick et al.\\ 2009).  In conclusion, photometric redshifts are useful tools for studying galaxy clusters. They enable efficient and complete pre-selection of cluster members for spectroscopy, allow accurate determinations of the cluster redshifts based on photometry alone, provide a means of determining cluster membership, especially for bright sources, and can be used to search for galaxy clusters."
    },
    "0910/0910.4905_arXiv.txt": {
        "abstract": "We use time-varying models of the coupled evolution of the HI, $\\rm H_2$ gas phases and stars in galaxy-sized numerical simulations to: a) test for the emergence of the Kennicutt-Schmidt (K-S) and the $\\rm H_2$-pressure relation, b) explore a realistic $\\rm  H_2$-regulated star formation recipe which brings forth a neglected and  potentially significant SF-regulating factor, and c) go beyond typical  galactic environments (for which these galactic empirical relations are  deduced) to explore the early evolution of very gas-rich galaxies.  In this work we model low mass galaxies ($M_{\\rm baryon} \\le 10^9 \\msun$), while incorporating an independent treatment of CO formation and destruction, the most important tracer molecule of H2 in galaxies, along with that for the H2 gas itself. We find that both the K-S and  the $\\rm H_2$-pressure  empirical relations  can robustly  emerge in galaxies after a dynamic equilibrium sets in between the various ISM states, the  stellar component and its feedback  ($\\rm T\\ga 1\\,Gyr$). The only  significant dependence of  these relations seems to  be for the  CO-derived (and  thus directly observable) ones,  which  show a strong dependance  on the  ISM metallicity.  The  $\\rm H_2$-regulated star formation recipe successfully reproduces the morphological and quantitative aspects of previous numerical models while doing away with the star formation efficiency parameter.  Most of the $\\rm HI\\rightarrow H_2$ mass  exchange is found taking place under highly    non-equilibrium    conditions necessitating    a time-dependent treatment even in  typical ISM environments. Our dynamic models indicate that the CO molecule can be a poor, non-linear, $\\rm H_2$  gas tracer.  Finally, for early evolutionary  stages ($\\rm T\\la 0.4\\,Gyr$) we find significant and systematic deviations of the true star formation from  that expected  from  the K-S  relation,  which are  especially pronounced and prolonged for  metal-poor systems.  The  largest such deviations  occur for the  very gas-rich galaxies, where deviations of  a factor $\\sim  3-4$ in global star formation rate  can take place with  respect to those  expected from the CO-derived  K-S relation.  This is  particularly important since gas rich systems at high redshifts  could  appear as having unusually high star-formation rates with respect  to their CO-bright $\\rm H_2$ gas reservoirs. This points to a possibly serious deficiency of  K-S relations  as  elements  of the  sub-grid  physics  of  star formation  in simulations  of   structure  formation  in  the  Early Universe. ",
        "introduction": "}  In spite of the fact that the general character  of the cycle through which galaxies  convert their ISM to  stars has been known  for a long time,  it  has  proven  to  be remarkably  difficult  to  formulate  a predictive theoretical framework  for this process. Indeed, we know that throughout most of the Universe star formation takes place in molecular gas complexes \\citep[e.g.][]{Solomon2005, Omont2007}. The $\\rm  HI\\rightarrow H_2$  phase  transition is conditioned by a combination  of sufficiently high HI column densities and pressures, and consequently star formation tends to concentrate in high  density regions,  e.g.  in  the  central parts  of galaxies,  in spiral arms or high-pressure concentrations  of gas formed by bulk gas motions or swept up by the shocks from supernovae and stellar winds of OB associations.  The link between molecular gas and star formation is so tight that  it has even been used to infer  the distribution of the former by that of the  latter, when the CO-$\\rm H_2$ conversion factor was still  considered  very uncertain  \\citep{Rana1986}.   In the  Galaxy, where   this  link   is  best   studied   \\citep[e.g.][and  references therein]{Blitz1997}, the latest  results confirm star formation always taking  place  in  CO-bright  molecular  clouds, even  at  very  large galactocentric  distances  \\citep{Kobayashi2008}.   In other  galaxies this tight association has been verified in all cases where sufficient angular resolution is  available \\citep[e.g.][]{Wong2002}.  Thus it is fair to say that {\\it  $\\rm H_2$ formation is a necessary prerequisite for star  formation in  galaxies,}  and incorporating  it in  galaxy-sized numerical simulations  of gas and  stars is the single  most important step currently missing from a realistic rendering of star formation in such models.  Following the early  and widespread observational evidence establishing the $\\rm H_2$-(star formation) link, the  inclusion of the $\\rm H_2$ gas phase in   numerical  models   of   galaxies  has   occured  only   recently \\citep{Pelupessy2006, Dobbs2006, Robertson2008}, and the refinement of these models is an area of ongoing research. This is mainly due to the difficulty of tracking the dynamic and thermodynamic evolution of $\\rm H_2$ and its precursor phase, the Cold Neutral Medium HI \\citep[$\\rm n \\sim 5-100$  cm$^{-3}$, $\\rm T_k  \\sim 60-200$K,][]{Wolfire2003} in galaxy-sized numerical models and due to the strong $\\rm H_2$ self-shielding complicating radiative transfer models of its far-UV radiation-induced destruction.   The  first  problem  has prevented  most efforts  from properly  tracking the  $\\rm  HI\\rightarrow H_2$  phase transition  in galaxies  without  resorting  to  simplifying steady-state  solutions \\citep{Hidaka2002,   Robertson2008} (suitable   only   for  quiescent galactic environments), or  to semi-empirical multiphase models \\citep[e.g.][]{Semelin2002} with limited predictive value.  The second problem confounds even  numerical simulations tracking  the $\\rm HI\\rightarrow H_2$ phase transition in individual gas  clouds where local approximations of the self-shielding HI/$\\rm H_2$  volume (necessary for  numerically manageable solutions) can make the  $\\rm H_2$ gas  mass fraction  a strong function       of      the      chosen       numerical      resolution \\citep[e.g.][]{Glover2007a}.   Finally   a  secondary,  yet  important problem of such single gas  cloud simulations is posed by the constant boundary  conditions  assumed during  their  evolution,  which are  an unlikely setting for  real gas clouds immersed in  the ISM environment of a galaxy.  In such environments cloud boundary conditions that powerfully  influence the  $\\rm HI\\rightarrow  H_2$  phase transition, such  as the  ambient  FUV  radiation field  and  pressure, change  on timescales comparable  or shorter than ``internal''  cloud dynamic and thermodynamic  timescales,  especially   in  vigorously  star  forming environments \\citep[e.g.][]{Parravano2003, Wolfire2003, Pelupessy2006}.   Despite  the   aforementioned  difficulties the incorporation of the  $\\rm H_2\\leftrightarrow HI$  gas phase interplay, and its strong role as star formation regulator, in numerical models holds the promise of large  improvements in  their handling  of galaxy  evolution,  and the possibility  of unveiling  new,  hitherto neglected,  aspects of  star formation  feedback  on   the  ISM.   In  this  paper   we  apply  our numerical  models for the  coupled evolution of  gas (HI, $\\rm H_2$)  and  stars  \\citep{Pelupessy2006}  to the evolution of low mass galaxies ($M_{\\rm baryon}<10^9 \\msun$) in order to explore two  new  key directions, the first of which is  the  emergence  of  two important  empirical relations   found   for   galaxies   in  the   local Universe:   the Kennicutt-Schmidt (K-S)  and the $\\rm H_2$--pressure  relation. Secondly we will make a investigation of very gas-rich systems (more typical of the Early Universe), and check whether the aforementioned  relations remain  valid during  their evolution.  The latter  is  of  crucial  importance  given  the  prominant  role  such empirical   relations   are   given   in   describing   the   sub-grid star-formation/gas  interplay in  cosmological  simulations of  galaxy evolution (where the resolution  limitations imposed by the simulation of large  volumes preclude a detailed description  of star formation). Finally along  with our original time-dependent treatment  of the $\\rm HI\\rightarrow H_2$  phase transition we also include the  CO molecule, allowing  direct   comparisons  to   the  {\\it  observed}   $\\rm H_2$  gas distributions, and a new independent investigation of  the   CO-$\\rm H_2$ relation  within  the  dynamical   setting  of  an evolving galaxy.  The  structure  of  our  work  is  as  follows:  in section~\\ref{sec:model}  we present  relevant features  of  the model, show  semi-analytical predictions for  the $\\rm H_2$-pressure  relation, and formulate  the  extention  of  the  model  so  that  includes  CO,  in section~\\ref{sec:sims} we present  our detailed numerical simulations, and  investigate  the  K-S,  $\\rm H_2$--pressure and  CO--$\\rm H_2$  empirical relations.  In  section~\\ref{sec:disc} we investigate  and discuss the validity of  the important K-S relation during  early galaxy evolution stages,  and  for  very   gas-rich  systems.  We  then  summarize  our conclusions in section~\\ref{sec:concl}. ",
        "conclusions": "\\label{sec:concl}  We use  our time-varying, galaxy-sized, numerical  models of gas+stars that track the ISM thermodynamics and the $\\rm HI\\leftrightarrow H_2$ gas phase exchange, to investigate: a) the emergence of two prominent empirical relations  deduced for galaxies in the  local Universe: the Kennicutt-Schmidt (K-S) relation  and the $\\rm H_2$-pressure relation, b) the  effects  of  a  more realistic  $\\rm H_2$-regulated  star  formation recipe, and c) the evolution of very gas-rich systems. Our models now include  a separate treatment  for formation  and destruction  of the $\\rm H_2$-tracing CO  molecule, which allows a direct  comparison of such models with observations, and  a new independent investigation of the CO-$\\rm H_2$ concomitance in the  ISM of evolving galaxies. Our findings can be summarized as follows  \\begin{itemize}  \\item  For  ISM states  of  $\\rm H_2$/HI  equilibrium, an  $\\rm H_2$-pressure relation close to  the one observed robustly emerges  for a wide range  of  parameters,  with   a  strong  dependance  mostly  on metallicity.  For  the more realistic  non-equilibrium $\\rm H_2$/HI states   {\\it  only   the   CO-bright  $\\rm H_2$   phase  shows   an $\\rm H_2$-pressure relation similar to the one observed.}  \\item The $\\rm H_2$-regulated star  formation model successfully models star formation without  the adhoc  parameter of the local star formation efficiency adopted   by   most   galaxy-sized  numerical   models,   while incorporating  a  fundamental  aspect  of  the  star  formation process.  \\item  A   comparison  between   numerical  models  using   the  usual simple-delay  (SD)  and the  new  molecular-regulated (MR)  star formation  recipes reveals  very  few differences.  It shows  a factor of $\\sim  3-4$ more efficient star formation  per CNM gas mass than the case of MR star formation.  \\item  We  find  little   sensitivity  of  the  global  SF  efficiency M(HI+$\\rm H_2$)/SFR   to  the  SF   recipe  chosen,   once  dynamic equilibrium  between  ISM  phases  and stars  is  established, yielding confidence  to (K-S)-type  of relations emerging  as a general characteristic of galaxies.   \\item A  non-equilibrium $\\rm HI\\leftrightarrow H_2$ gas  mass exchange is revealed  taking   place    under   typical   ISM   conditions, demonstrating the need for a full dynamic rather than stationary treatment of these ISM  phases.   \\item The  CO molecule can be  a poor, non-linear, tracer  of the true underlying  $\\rm H_2$  gas  distribution, especially  in  metal-poor systems, and  even in  those with very  high gas  mass fractions (more typically found at high redshifts).   \\item A K-S relation  robustly emerges from our time-dependent models, irrespective  of   the  SF   recipe  used,  after   a  dynamical equilibrium is established  ($\\rm T\\ga $1\\,Gyr). The CO-derived K-S relation has a more shallow slope than the one involving the total gas mass, and as  in the $\\rm H_2$-pressure relation, a strong dependance on metallicity is found.   \\item  At  early evolutionary  timescales  ($\\rm  T\\la 0.4\\,Gyr$)  our models show  {\\it significant and systematic deviations  of the true star formation from that expected from the K-S relation,} which seem especially pronounced  and prolonged for  metal-poor systems.  These deviations  occur even  for the  CO-derived K-S  relation  (the more realistic one since CO is directly observable and traces the densest $\\rm H_2$ gas  which ``fuels'' star formation), and  even for metal-rich systems where CO tracks the $\\rm H_2$ gas well.   \\item  The largest  deviations  from  the K-S  relation  occur at  the earliest evolutionary stages of  the systems modeled here ($\\rm T\\la 0.2 Gyr$) and for the most gas-rich ones. During this time significantly higher  star formation rates  per CO-bright $\\rm H_2$ gas mass  occur, and such  star-forming galaxies may  have been already observed at high redshifts.  \\end{itemize}  Finally  we must  note that  when it  comes to  the  gas-rich galaxies accessible to current  observational capabilities at high redshifts, our   results,  drawn   for  much   less  massive   systems,  remain provisional.   Nevertheless for  more massive  gas-rich  systems the larger amplitudes of  ISM equilibrium-perturbing agents (e.g.  SNs, far-UV  radiation  fields), and  the  shorter  timescales that  will characterize their variations are more likely than not to exaggerate the deviations  of true star  formation versus the one  derived from (K-S)-type  phenomenological relations.   A  dedicated observational effort  to  study  such  galaxies  at high  redshifts  (soon  to  be dramatically  enhanced  by  ALMA),  as well  as  extending  detailed numerical  modeling   of  gas  and  stars  to   larger  systems  (as computational  capabilities  improve),  can help  establish  whether (K-S)-type relations  remain valid during  most of the  stellar mass built-up in galaxies, or  only emerge after dynamic equilibrium has been reached during much latter evolutionary stages."
    },
    "0910/0910.2242_arXiv.txt": {
        "abstract": "We present a spectroscopic sample of 910 distant halo stars from the Hypervelocity Star survey from which we derive the velocity dispersion profile of the Milky Way halo.  The sample is a mix of 74\\% evolved horizontal branch stars and 26\\% blue stragglers.  We estimate distances to the stars using observed colors, metallicities, and stellar evolution tracks.  Our sample contains twice as many objects with $R>50$ kpc as previous surveys.  We compute the velocity dispersion profile in two ways:  with a parametric method based on a Milky Way potential model, and with a non-parametric method based on the caustic technique originally developed to measure galaxy cluster mass profiles.  The resulting velocity dispersion profiles are remarkably consistent with those found by two independent surveys based on other stellar populations:  the Milky Way halo exhibits a mean decline in radial velocity dispersion of $-0.38\\pm0.12$ \\kms\\ kpc$^{-1}$ over $15<R<75$ kpc.  This measurement is a useful basis for calculating the total mass and mass distribution of the Milky Way halo. ",
        "introduction": "A fundamental property of the Milky Way galaxy is its total mass.  In the Cold Dark Matter paradigm, the total mass of a galaxy's dark matter halo correlates with its merger history, star formation history, and number of its satellite sub-halos.  A galaxy's mass distribution is also a fundamental constraint on theories.  Cold dark matter models predict that density follows an universal NFW profile \\citep{nfw97}; in Modified Newtonian Dynamics there is no dark matter halo and the mass distribution is highly flattened compared to cold dark matter models \\citep{hernandez09}.  The Milky Way provides an unique laboratory for testing these issues.  Nevertheless, the mass and mass distribution of the Milky Way are among the most poorly known of all Galactic parameters, especially at large distances $R \\gtrsim 50$ kpc.  Total mass estimates span at least a factor of 4, from $0.5\\times10^{12}$ to $2\\times10^{12}$ M$_{\\odot}$ (see below).  Recent theoretical work based on semi-analytic models and timing arguments suggest the Milky Way may be even more massive \\citep{li08, abadi09}.  Clearly, the mass and mass distribution of the Milky Way remain controversial and important issues.  A powerful approach for measuring the Milky Way's mass distribution is measuring the motions of tracer objects.  Radio masers probe the Galactic rotation curve out to 15 kpc \\citep{reid09}; H {\\sc i} gas probes the rotation curve to larger distances \\citep{kalberla08}.  Luminous tracers at $R \\gtrsim 50$ kpc, however, are rare. Historically, globular clusters and dwarf galaxies were used to measure the total mass of the Milky Way \\citep[e.g.][]{little87, zaritsky89, kulessa92, kochanek96, wilkinson99}.  These mass estimates were based on samples of a few dozen objects, and suffered from a systematic factor of $\\sim$2 uncertainty depending on the inclusion of Leo I.  Post-main sequence halo stars provide denser tracers, but the blue horizontal branch (BHB) star samples of \\citet{sommer97} and \\citet{sakamoto03} are limited to $R\\lesssim20$ kpc.  \\citet{battaglia05} added 58 distant red giants from the Spaghetti survey and claimed a large decline in the velocity dispersion at $R>$50 kpc.  Recently, \\citet{xue08} analyzed the Sloan Digital Sky Survey (SDSS) sample of 2,466 BHB stars and found a small decline in velocity dispersion with distance.  Only 80 of the SDSS BHB stars are located at $R>50$ kpc.  Here we use the distant halo stars from the hypervelocity star (HVS) program \\citep{brown05, brown06, brown06b, brown07a, brown07b, brown09a, brown09b} to measure the velocity dispersion profile of the Milky Way.  Our dataset is a complete spectroscopic sample of 910 stars observed over 7300 deg$^2$ of the SDSS Data Release 6 imaging region.  We gain increased leverage on the velocity dispersion profile for $R \\gtrsim 50$ kpc.  We find an average decline in radial velocity dispersion of $0.38\\pm0.12$ \\kms\\ kpc$^{-1}$ over $15<R<75$ kpc.  In \\S 2 we describe the observations and luminosity estimates for stars in our sample.  In \\S 3 we make parametric and non-parametric estimates of the velocity dispersion profile.  We also discuss possible systematics, and compare our results with earlier work.  We conclude in \\S 4.  The data are in the Appendix. ",
        "conclusions": "The mass and mass distribution of the Milky Way are fundamental parameters because they link directly to theoretical models. We use a spectroscopic sample of 910 halo stars derived from our HVS survey to measure the velocity dispersion profile of the Milky Way.  The stars are 74\\% BHB stars and 26\\% blue stragglers.  We estimate luminosities using stellar evolution tracks for metal poor main sequence stars and BHB stars.  Because of the non-Gaussian distribution of luminosity estimates, we use a Monte Carlo approach to calculate velocity dispersion and its uncertainty.  We calculate the velocity dispersion profile in two ways:  a parametric method based on a $v_{esc}(R)$ model, and a non-parametric method based on the caustic technique originally developed to measure galaxy cluster mass profiles. Comparing the two methods provides a measure of the systematic uncertainty arising from the clipping of outliers in velocity. The velocity dispersion from the caustic method is 10\\% smaller than the $v_{esc}(R)$ method, but both methods identify a similar decline in velocity dispersion with distance:  $-0.38\\pm0.12$ \\kms\\ kpc$^{-1}$, valid for $15<R<75$ kpc.  Our sample contains a factor of 8 more stars than \\citet{battaglia05} and a factor of 2 more stars than \\citet{xue08} at $R>50$ kpc.  The velocity dispersion profiles observed by these independent datasets are consistent at the 1.5-$\\sigma$ level, and have an average velocity dispersion slope identical to our result. Remarkably, no matter what tracers are used, observers find the same halo velocity dispersion profile.  The velocity dispersion profile is a basis for measuring the total mass and mass distribution of the Milky Way halo. A companion paper by Gnedin et al.\\ (in preparation) presents the theoretical calculations that turn the observed velocity dispersion profile into a mass determination of the Milky Way.  For further progress in measuring the Milky Way velocity dispersion profile, it is essential to identify tracers at distances $R>50$ kpc. It is difficult to find $R>50$ kpc tracers because of the steep decline in the density of the stellar halo.  It is also very difficult for proper motion surveys to measure tracers at $R>50$ kpc distances. On-going spectroscopic radial velocity surveys, such as the SDSS-3 survey and our own HVS survey, promise to better trace the Milky Way in coming years."
    },
    "0910/0910.5654_arXiv.txt": {
        "abstract": "We have invented a novel technique to measure the radio image of a pulsar scattered by the interstellar plasma with 0.1~mas resolution.   We extend the ``secondary spectrum'' analysis of parabolic arcs by Stinebring et al.\\ (2001) to very long baseline interferometry and, when the scattering is anisotropic, we are able to map the scattered brightness astrometrically with much higher resolution than the diffractive limit of the interferometer.  We employ this technique to measure an extremely anisotropic scattered image of the pulsar B0834$+$06 at 327~MHz. We find that the scattering occurs in a compact region about 420~pc from the Earth. This image has two components, both essentially linear and nearly parallel. The primary feature, which is about 16~AU long and less than 0.5~AU in width, is highly inhomogeneous on spatial scales as small as 0.05~AU. The second feature is much fainter and is displaced from the axis of the primary feature by about 9~AU. We find that the velocity of the scattering plasma is $16 \\pm 10$~\\kms approximately parallel to the axis of the linear feature. The origin of the observed anisotropy is unclear and we discuss two very different models. It could be, as has been assumed in earlier work, that the turbulence on spatial scales of  ($\\sim 1000$~km) is homogeneous but anisotropic. However it may be that the turbulence on these scales is homogeneous and isotropic but the anisotropy is produced by highly elongated (filamentary) inhomogeneities of scale 0.05-16 AU. ",
        "introduction": "Radio pulsars provide a powerful tool for studying the ionized interstellar medium.  The dispersion in their pulse arrival times probes the mean electron density.  Their very small diameters ensure that they display the full range of scintillation and scattering phenomena, which probe the fine spatial structure in the electron density.  Their pulse amplitudes exhibit a combination of diffractive and refractive intensity scintillation on times from seconds to months.   Compact emission from some bright active galactic nuclei can also show scintillation on times of hours to months, albeit smoothed by the effect of their larger angular diameters.   The large body of pulsar scintillation data has been interpreted in terms of homogeneous isotropic Kolmogorov turbulence in the interstellar plasma \\citep{ars95}.  See reviews by \\citet{Ric90,Nar92}.  However, a litany of observational evidence now points to the existence of compact ionized structures in the interstellar medium (ISM) whose scattering characteristics are well beyond those of such homogeneous isotropic Kolmogorov turbulence.  In particular the scattering is seldom uniformly distributed along the line of sight. It is often dominated by one local region somewhere in the line of sight, which we refer to as a ``thin screen.''  Although it is unlikely to resemble a screen, it is thin with respect to the total line of sight from the source to the observer \\citep[e.g.,][]{putney,rjtr}.   Inhomogeneity in the turbulence is required on kiloparsec scales to explain how the level of pulsar scattering varies with distance and Galactic coordinates \\citep{Cor91}. Inhomogeneity is also required on the parsec scale in the local ISM to explain the intermittent nature of the scintillation observed in a few quasars on hour-long time scales \\citep[e.g.,][]{DTdB03,KC06}.  In addition, evidence for AU-scale inhomogeneity in the turbulence comes from extreme scattering events (ESEs), which are observed in a few quasars as rare, large (10 to 50\\%) variations in flux density over several weeks \\citep{Fiedler87,La01,Sen08}.  These are generally seen as a decrease in flux density attributed to the passage across the line of sight of an ionized cloud which either scatters \\citep{Fiedler87} or refracts \\citep{Rom87} the radiation.  In recent years there has also been increasing evidence that interstellar scattering (ISS) is not only inhomogeneous but also anisotropic.  The most direct measure of anisotropy is through very long baseline interferometry (VLBI) imaging of scattered brightness, but it is only detectable on a few heavily scattered lines of sight \\citep[e.g.,][]{Des01}.  Anisotropy in the ISS diffraction pattern has also been measured indirectly \\citep[e.g.,][]{Ric02, DTdB03,Col05,Big06}.  Such observations give evidence for elongated fine structure in the ISM on scales of thousands of kilometers, suggesting anisotropic magneto-hydrodynamic turbulence controlled by the magnetic field as discussed by \\citet{Gol95} and \\citet{Spa99}.  The examples cited above show departures from both homogeneity and isotropy in the ionized ISM. It is possible, but by no means proven, that the various phenomena have a common origin in a population of AU-scale anisotropic regions of enhanced density and turbulence which we here generically refer to as ``clouds.''  However, we note that such localized clouds must contain fine scale substructure that causes scattering at radio frequencies.  Such a population of clouds presents a serious puzzle.  Their number density must be many orders of magnitude greater than that of stars and their implied electron densities $n_{\\rm e} \\age 10$~cm$^{-3}$ are much higher than expected in pressure equilibrium in the warm ionized phase ISM \\citep{Rom87}.  The discovery of parabolic arcs in the ISS of pulsars by \\citet{Sti01} has added a powerful new tool for probing such clouds.  Many pulsars exhibit parabolic arcs in their secondary spectra (SS), which is the power spectrum (versus delay and Doppler frequency) of the dynamic spectrum of intensity (versus frequency and time).  The arcs are sometimes narrow (in delay) which implies scattering by a thin layer.   The distribution of power in the SS often reveals anisotropic scattering and in some cases there are discrete downwards facing ``arclets,'' which also imply scattering from isolated anisotropic clouds (``cloudlets''). See \\citet{Cor06,Wal04} for interpretation of the arcs.  The SS allows a two dimensional reconstruction of the scattered image from observations at a single receiver, since each point in the SS isolates the scintillation power associated with interference between pairs of points on the scattering disk.  However the reconstruction can be model dependent and has an inherent two-fold ambiguity \\citep{Cor06,Tra07}.  Occasionally one can also see isolated peaks in the SS corresponding to narrow bandwidth fringes in the dynamic spectrum \\citep{Wol87,Ric97}.  Such fringes result from interference between the normal primary (on-axis) scattering disk and the off-axis discrete cloud.  Recent results from \\citet{Hil05} have compounded the difficulties in understanding the nature of the underlying ionized clouds.  They observed the SS towards pulsar B0834$+$06 and found four distinct arclets scattered through 7 to 12~mas which they interpreted as originating from $\\sim$0.2~AU clouds requiring $n_{\\rm e} > 100$~cm$^{-3}$, similar to those invoked to explain ESEs towards quasars.  By monitoring the evolution of structures in the secondary spectrum, they followed these clouds over three weeks and showed that they co-moved with the rest of the scattering material.  In this paper we report VLBI observations of the scintillation from the same pulsar (B0834$+$06) in order to further investigate these clouds.  We have developed a novel astrometry technique that makes use of SS-like quantities derived from the interferometer visibilities.  Using these ``secondary cross spectra'' (defined in Table~\\ref{tab:defs}), we can accurately localize points on the scattering screen corresponding to high signal-to-noise pixels of the SS. We use the results of the astrometry of many such points to measure the distance and velocity of the interstellar clouds and so define a precise model for the scattering.  Applying this precise model to the astrometric image with the scattering model allows us to eliminate all the otherwise troublesome ambiguities and validates the model.  We then use this precise model to recover the scattered image with even greater angular resolution from the secondary spectrum itself.  In our observations the Rayleigh resolution (i.e., synthesized beam) of the VLBI array is about 35~mas, the astrometric precision is about 1~mas, and structure in the scattered image recovered by modeling is found on a scale of $100$~$\\mu$as. ",
        "conclusions": "This paper describes a novel VLBI technique resulting in a two-dimensional image of the scattering screen of pulsar B0834$+$06. The baseband data that were recorded allowed high resolution dynamic spectra to be produced. The secondary spectra produced with the two-hour dynamic spectra could allow sharply defined arclets to be identified with delays as high as 1~ms.   The scattered image was developed by astrometrically mapping points chosen from the secondary spectrum to bright points in the sky plane. These points were clustered in two clearly defined groups: a primary scattering disk which is elongated and inclined $27 \\pm 2^{\\circ}$ to the pulsar proper motion direction and a second, non-colinear, feature corresponding to the 1~ms feature of the secondary spectrum.  Diagnostic measurements place the two features at essentially the same distance, 65\\% of the way to the pulsar. The two-dimensional distribution of points in the scattered image allows both transverse components of the effective velocity to be determined via relationships connecting the Doppler frequency with location in the image.  The discrete feature at 1~ms delay contains about 4\\% of the total received power. This feature is expected to to be visible only for a few weeks during which time its delay should drift as the pulsar moves; the impact on timing this pulsar at $\\sim 327$~MHz due to such a feature is a time variable wander with magnitude $\\sim 40$~$\\mu$s. This should come as a caution to those aiming to perform precision pulsar timing at low frequencies on pulsars that exhibit the extreme forms of scintillation that are characteristic of B0834$+$06. Further, pulsars with sub-microsecond structure may experience apparent pulse profile evolution yielding additional complications in their timing.   We were able to estimate the effective scintillation velocity vector, which depends on a distance-weigthed sum of the velocities of the pulsar, the Earth and the sacttering plasma.  By using the published proper motion we estimated the velocity of the scattering plasma to be $16\\pm 10$~\\kms approximately parallel to the scattering axis.  Since the errors in this interesting result are dominated by the uncertainty in the pulsar proper motion, we have undertaken a new set of VLBI measurements to improve its precision.  The interpretation of ISS in pulsars has often assumed isotropy in the scattering.  The extremely anisotropic scattering found here would substantially alter any quantitative modelling of the plasma were it to be a common feature in other regions on the interstellar medium.  A description of the underlying plasma physics must await a resolution of the two possible geometries mentioned in the previous section, but the results and the method provide an exciting new glimpse of the ionized ISM at scales of 0.1 to 10~AU."
    },
    "0910/0910.0071_arXiv.txt": {
        "abstract": "{In mean-field magnetohydrodynamics the mean electromotive force due to velocity and magnetic field fluctuations plays a crucial role. In general it consists of two parts, one independent of and another one proportional to the mean magnetic field. The first part may be nonzero only in the presence of mhd turbulence, maintained, e.g., by small-scale dynamo action. It corresponds to a battery, which lets a mean magnetic field grow from zero to a finite value. The second part, which covers, e.g., the $\\alpha$ effect, is important for large-scale dynamos. Only a few examples of the aforementioned first part of mean electromotive force have been discussed so far. It is shown that a mean electromotive force proportional to the mean fluid velocity, but independent of the mean magnetic field, may occur in an originally homogeneous isotropic mhd turbulence if there are nonzero correlations of velocity and electric current fluctuations or, what is equivalent, of vorticity and magnetic field fluctuations. This goes beyond the Yoshizawa effect, which consists in the occurrence of mean electromotive forces proportional to the mean vorticity or to the angular velocity defining the Coriolis force in a rotating frame and depends on the cross-helicity defined by the velocity and magnetic field fluctuations. Contributions to the mean electromotive force due to inhomogeneity of the turbulence are also considered. Possible consequences of the above and related findings for the generation of magnetic fields in cosmic bodies are discussed.} ",
        "introduction": "Mean--field magnetohydrodynamics has proved to be a useful tool for studying the behavior of mean magnetic fields in turbulently moving electrically conducting fluids (see, e.g., Moffatt 1979, Krause \\& R\\\"adler 1980, Brandenburg \\& Subramanian 2005). Within this framework both the magnetic field $\\bB$ and the fluid velocity $\\bU$ are split into mean parts, $\\bmB$ and $\\bmU$, and fluctuating parts, $\\bb$ and $\\bu$. Starting from the induction equation governing $\\bB$ it is concluded that the mean magnetic field $\\bmB$ has to obey \\EQ \\p_t \\bmB = \\eta \\bnab^2 \\bmB + \\bnab \\x (\\bmU \\x \\bmB + \\bscE ) \\, , \\quad \\bnab \\cdot \\bmB = 0 \\, . \\label{eq01} \\EN Here, $\\eta$ means the magnetic diffusivity of the fluid, for simplicity considered as independent of position, and $\\bscE$ the mean electromotive force caused by the velocity and magnetic fluctuations, \\EQ \\bscE = \\lan \\bu \\x \\bb \\rangle \\, . \\label{eq03} \\EN Mean fields are defined by some kind of averaging satisfying the Reynolds rules. They are denoted either by overbars or synonymously by angle brackets.  The induction equation governing $\\bB$ also implies \\EQA && \\!\\!\\!\\!\\!\\!\\!\\!\\!\\!\\! \\p_t \\bb = \\eta \\bnab^2 \\bb + \\bnab \\x [ (\\bu \\x \\bb)' + \\bmU \\x \\bb + \\bu \\x \\bmB], \\nonumber\\\\ && \\qquad \\qquad \\qquad \\qquad \\qquad \\bnab \\cdot \\bb = 0 \\, , \\label{eq05} \\ENA where $(\\bu \\x \\bb)' = \\bu \\x \\bb - \\lan \\bu \\x \\bb \\ran$. With this in mind we may conclude that $\\bscE$ can be represented as a sum of two parts, \\EQ \\bscE = \\bscE^{(0)} + \\bscE^{(\\mB)} \\, , \\label{eq07} \\EN where $\\bscE^{(0)}$ is a functional of $\\bu$ and $\\bmU$, and $\\bscE^{(\\mB)}$ a functional of $\\bu$, $\\bmU$ and $\\bmB$, which is linear in $\\bmB$ but vanishes if $\\bmB$ is zero everywhere and at all past times (see, e.g., R\\\"adler 1976, 2000, R\\\"adler \\& Rheinhardt 2007). These statements apply independently on whether or not $\\bu$ or $\\bmU$ depend on $\\bmB$. If they depend on $\\bmB$ and the total variation of $\\bscE$ with $\\bmB$ is considered, $\\bscE^{(0)}$ may well depend on $\\bmB$, and $\\bscE^{(\\mB)}$ need not be linear in $\\bmB$.  A non--zero $\\bscE^{(0)}$ corresponds to a battery. Assume for a moment that equation (\\ref{eq01}) for $\\bmB$ with $\\bscE = \\bzo$ has no growing solutions. If then $\\bscE^{(0)}$ takes non--zero values, but $\\bscE^{(\\mB)}$ remains equal to zero, $\\bmB$ grows, even if initially equal to zero, to a finite magnitude determined by $\\bscE^{(0)}$. If, on the other hand, $\\bscE^{(0)}$ remains equal to zero a non--zero $\\bscE^{(\\mB)}$ may allow (if it has a suitable structure) a dynamo, that is, let an arbitrarily small seed magnetic field $\\bmB$ grow exponentially (in the absence of back--reaction on the fluid motion even endlessly). A small non--zero $\\bscE^{(0)}$ may deliver a seed field for such a dynamo. This possibility has been already discussed in the context of young galaxies (Brandenburg \\& Urpin 1998).  In most of the general representations and applications of mean--field magnetohydrodynamics the part $\\bscE^{(0)}$ of the electromotive force $\\bscE$ has been ignored. Indeed, if it occurs at all, it decays to zero in the course of time except in cases in which an independent magnetohydrodynamic turbulence exists, e.g., as a result of a small--scale dynamo.  The possibility of a non--zero $\\bscE^{(0)}$ due to local, that is small--scale, dynamos in the solar convection zone has been discussed by R\\\"adler (1976). We express his statements here by \\EQ \\bscE^{(0)} = c_\\gamma \\bgamma + c_\\Omega \\bOmega + c_{\\gamma \\Omega} \\bgamma \\x \\bOmega \\, , \\label{eq09} \\EN where $\\bgamma$ is a gradient, e.g., of the turbulence intensity, $\\bOmega$ the angular velocity responsible for the Coriolis force, and $c_\\gamma$, $c_\\Omega$ and $c_{\\gamma \\Omega}$ some coefficients. More precisely, $c_\\gamma$ and $c_{\\gamma \\Omega}$ are scalars and $c_\\Omega$ is a pseudoscalar.  Another interesting result has been derived by Yoshizawa (1990, see also Yoshizawa 1993 or  Yoshizawa, Itoh \\& Itoh 2003). Considering an originally homogeneous isotropic magnetohydrodynamic turbulence under the influence of a mean flow or a rigid--body rotation, or both, he found \\EQ \\bscE^{(0)} = c_W \\bW + c_\\Omega \\bOmega \\, , \\label{eq11} \\EN where $\\bW = \\bnab \\x \\bmU$ is the mean vorticity, $\\bOmega$ again the angular velocity responsible for the Coriolis force, and $c_W$ and $c_\\Omega$ are pseudoscalars coefficients which are, roughly speaking, proportional to the cross--helicity $\\lan \\bu \\cdot \\bb \\ran$. This result has recently been used for an interpretation of the Archontis dynamo (Sur \\& Brandenburg 2009).  The main purpose of this paper is to demonstrate that a mean electromotive force $\\bscE^{(0)}$ proportional to the mean fluid velocity $\\bmU$ may occur in originally homogeneous isotropic magnetohydrodynamic turbulence. This should be expected as soon as there is a non--zero correlation between the fluctuating parts of velocity and electric current, $\\bu$ and $\\bj = \\mu_0^{-1} \\bnab \\x \\bb$, or, what is equivalent, between the fluctuating parts of vorticity and magnetic field, $\\bomega = \\bnab \\x \\bu$ and $\\bb$; as usual, $\\mu_0$ means the magnetic permeability. We express this condition roughly by saying that $\\lan \\bu \\cdot \\bj \\ran$ or $\\lan \\bomega \\cdot \\bb \\ran$ have to be unequal to zero. Unlike $\\lan \\bu \\cdot \\bb \\ran$, which characterizes the linkage between vortex tubes and magnetic flux tubes, $\\lan \\bu \\cdot \\bj \\ran$ quantifies the linkage between vortex tubes and current tubes.  In Section \\ref{sec2} we explain the basis of our calculations and provide general relations for the determination of the mean electromotive force $\\bscE^{(0)}$. In Section \\ref{sec3} we derive results for homogeneous isotropic turbulence, in particular the last--mentioned one, and we also reproduce that given by (\\ref{eq11}). Proceeding then in Section \\ref{sec4} to inhomogeneous turbulence we report on results related to those indicated in (\\ref{eq09}). The relevance of the results obtained in this paper and the need of further work are discussed in Section \\ref{sec5}. ",
        "conclusions": "\\label{sec5}  The most remarkable result of our calculations is that the mean electromotive force $\\bscE^{(0)}$ in a homogeneous isotropic magnetohydrodynamic turbulence may have a contribution proportional to the mean fluid velocity $\\bmU$, that is, $\\bscE^{(0)} = c_U \\, \\bmU + \\cdots$. We have labeled the occurrence of this contribution as $\\lan \\bu \\cdot \\bj \\ran$ effect. The coefficient $c_U$ turned out to be in general unequal to zero if only a non--zero correlation exists between the fluctuating parts of the fluid velocity and the electric current density, $\\bu$ and $\\bj = \\mu_0^{-1} \\bnab \\x \\bb$, or between the fluctuating parts of the vorticity and the magnetic field, $\\bomega = \\bnab \\x \\bu$ and $\\bb$. As far as the second--order correlation approximation applies, $c_U$ vanishes for $\\eta = \\nu$, that is, it changes its sign if $\\nu / \\eta$ varies and passes through $\\nu / \\eta = 1$.  The occurrence of magnetohydrodynamic turbulence does not automatically imply non--zero correlations of $\\bu$ and $\\bj$, or $\\bomega$ and $\\bb$. It depends on the special circumstances whether, e.g.,  $\\lan \\bu \\cdot \\bj \\ran$ or $\\lan \\bomega \\cdot \\bb \\ran$ are different from zero and what their signs are. For this and other reasons more work is needed to explore the importance of the $\\lan \\bu \\cdot \\bj \\ran$ effect in specific settings. In general the $\\lan \\bu \\cdot \\bj \\ran$ effect is accompanied by the $\\lan \\bu \\cdot \\bb \\ran$ effects. It has then also to be investigated which of these effects dominates.  At first glance there seems to be a inconsistency of our result concerning the $\\lan \\bu \\cdot \\bj \\ran$ effect in so far as $\\bscE^{(0)}$ should not depend on the choice of the frame of reference but $\\bmU$ obviously does. We must however keep in mind that we have fixed the frame of reference in our calculation by assuming that there isotropic turbulence occurs in the limit $\\bmU \\to \\bzo$. When estimating the $\\lan \\bu \\cdot \\bj \\ran$ effect we have therefore to specify $\\bmU$ as the mean velocity of the fluid relative to the frame in which the assumed causes of turbulence (in simulations the forcing) would, in the absence of mean motion, just lead to isotropic turbulence.  There is still another issue which has to be considered when applying our result concerning the $\\lan \\bu \\cdot \\bj \\ran$ effect to a specific situation. The deviation of the turbulence from isotropy due to the homogeneous part of the mean motion is crucial for that effect. It has to be scrutinized whether such a deviation indeed occurs under the considered circumstances. The sometimes assumed ``Galilean invariance\" of the turbulence (e.g., Sridhar \\& Subramanian 2009), that is, its independence of that part of the mean motion, would exclude the $\\lan \\bu \\cdot \\bj \\ran$ effect.  In principle we could have determined $\\bscE^{(0)}$ in a frame in which $\\bmU$ vanishes. An anisotropy of the turbulence, as calculated in the frame used above, would lead to a non--vanishing contribution $\\bscE^{(00)}$ instead of $c_U \\bmU$, which then has to be considered as another description of the $\\lan \\bu \\cdot \\bj \\ran$ effect."
    },
    "0910/0910.0840_arXiv.txt": {
        "abstract": "We report the results of a study of the rest-frame ultraviolet (UV) spectrum of the Cosmic Eye (J213512.73$-$010143), a luminous ($L \\sim 2 L^\\ast$) Lyman break galaxy at $z_{\\rm  sys}=3.07331$ magnified by a factor of $\\sim 25$ via gravitational lensing by foreground mass concentrations at $z = 0.73$ and 0.33. The spectrum, recorded at high resolution and signal-to-noise ratio with the ESI spectrograph on the Keck\\,{\\sc ii} telescope, is rich in absorption features from the gas and massive stars in this galaxy. The interstellar absorption lines are resolved into two components of approximately equal strength and each spanning several hundred \\kms\\ in velocity. One component has a net blueshift of $-70$ \\kms\\ relative to the stars and H\\,{\\sc ii} regions and presumably arises in a galaxy-scale outflow similar to those seen in most star-forming galaxies at $z = 2$--3. The other is more unusual in showing a mean \\textit{redshift} of $+350$ \\kms\\  relative to $z_{\\rm sys}$; possible interpretations include a merging clump, or material ejected by a previous star formation episode and now falling back onto the galaxy, or more simply a chance alignment with a foreground galaxy. In the metal absorption lines, both components  only partially cover the OB stars against which they are being viewed. However,  there must also be more pervasive diffuse gas to account for the near-total covering fraction of the strong damped \\lya\\ line, indicative of a column density $N$(H\\,{\\sc i})\\,$ = (3.0 \\pm 0.8) \\times 10^{21}$\\,cm$^{-2}$. We tentatively associate this neutral gas with the redshifted component, and propose that it provides the dust `foreground screen' responsible for the low ratio of far-infrared to UV luminosities of the Cosmic Eye.  The C\\,{\\sc iv} P~Cygni line in the stellar spectrum is consistent with continuous star formation with a Salpeter initial mass function, stellar masses from 5 to 100\\,\\msun, and a metallicity $Z \\sim 0.4 Z_{\\odot}$. Compared to other well-studied examples of strongly lensed galaxies, we find that the young stellar population of the Cosmic Eye is essentially indistinguishable from those of the Cosmic Horseshoe and MS\\,1512-cB58. On the other hand, the interstellar spectra of all three galaxies are markedly different, attesting to the real complexity of the interplay between starbursts and ambient interstellar matter in young galaxies observed during the epoch when cosmic star formation was at its peak. ",
        "introduction": "\\label{sec:introduction}  High redshift star-forming galaxies are often identified by a break in their ultraviolet continuum that is due to the Lyman limit, partly from interstellar (within the galaxy) and primarily from intergalactic H\\,{\\sc i} absorption below 912\\,\\AA\\ (Steidel et al. 1996).  Since the discovery of these `Lyman break galaxies' (or LBGs), samples of galaxies at $z = 2$--4 have increased a thousand-fold. Despite the variety of methods now employed to select high redshift galaxies, LBGs remain the most common and most extensively studied class of such objects.  Detailed studies of the physical properties of \\textit{individual} galaxies have in general been limited by their faintness ($m_{\\cal R}^{\\ast} = 24.4$, Steidel et al. 1999; Reddy et al. 2008).  In a few cases, however, gravitational lensing by foreground massive galaxies, groups, or clusters has afforded rare insights into the stellar populations and interstellar media of galaxies at $z = 2$--4 (e.g. Pettini et al. 2000, 2002; Teplitz et al. 2000; Lemoine-Busserolle et al. 2003; Smail et al. 2007; Swinbank et al. 2007; Cabanac, Valls-Gabaud, \\& Lidman 2008; Siana et al. 2008, 2009; Finkelstein et al. 2009; Hainline et al. 2009; Quider et al. 2009; Swinbank et al. 2009; Yuan \\& Kewley 2009 and references therein). The order-of-magnitude boost in flux provided by strong lensing makes it possible to record the spectra of these galaxies with a combination of high resolution and high signal-to-noise ratio (S/N) which for normal, unlensed galaxies will have to wait until the next generation of 30+\\,m optical/infrared telescopes.  Focusing on the rest-frame UV spectra in particular, our group has so far published observations of two of these strongly lensed high redshift galaxies: MS\\,1512-cB58 and J1148+1930 (cB58 and the `Cosmic Horseshoe' respectively; Pettini et al. 2002; Quider et al. 2009).  Both are $\\sim L^{*}$ galaxies at redshifts $z\\sim 2.5$ magnified by a factor of $\\sim 30$ (Seitz et al. 1998; Dye et al. 2008).  A considerable amount of data that are inaccessible to typical low-resolution studies is contained in these rest-UV spectra, including information on interstellar gas composition and kinematics, on the initial mass function (IMF) of starbursts at high $z$, and clues to the geometry of the gas and stars.  One of the motivations for this work is to assess the range of properties possessed by high redshift star-forming galaxies. By comparing the UV spectra of cB58 and the Cosmic Horseshoe, Quider et al. (2009) showed that the young stellar populations of these galaxies are very similar, as are the metallicities and kinematics of their interstellar media. On the other hand, the two galaxies exhibit clear differences in the covering fractions of their stars by the interstellar gas, differences which are reflected in the strikingly different morphologies of their Lyman~$\\alpha$ lines (a damped absorption profile in cB58 and a double-peaked emission profile in the Horseshoe). Studies at other wavelengths (Teplitz et al. 2000; Hainline et al. 2009) have also shown that cB58 and the Cosmic Horseshoe have comparable star formation rates, $\\rm{SFR} \\sim 50$--100\\,\\msun~yr$^{-1}$, at the upper end of the range of values measured in star-forming galaxies at $z = 2$--3, and similar dynamical masses, $M_{\\rm dyn} \\sim 1 \\times 10^{10}$\\,\\msun, typical of luminous galaxies at this epoch (Pettini et al. 2001; Erb et al. 2006b,c;  Reddy et al. 2006).  Clearly, several examples of strongly lensed galaxies need to be studied at such levels of detail for a full characterization of galaxy properties at these redshifts. Fortunately, many new cases have been discovered recently, mostly from dedicated searches in imaging data from the Sloan Digital Sky Survey (Estrada et al. 2007; Belokurov et al. 2007, 2009; Ofek et al. 2008; Shin et al. 2008; Kubo et al. 2009; Lin et al. 2009; Wen et al. 2009). Other surveys have capitalised on the superior spatial resolution of \\textit{Hubble Space Telescope (HST)} to identify high redshift galaxies strongly lensed into multiple images or arcs; the $z = 3.0733$ galaxy which is the focus of the present paper, dubbed the `Cosmic Eye', was indeed discovered by Smail et al. (2007) from an \\textit{HST} Snapshot program targeting high luminosity X-ray clusters.  In this paper we present high resolution observations of the rest-frame UV spectrum of the Cosmic Eye. The paper is organized as follows. In Section~\\ref{sec:eyesummary} we summarize the known properties of the Cosmic Eye from previous studies at a variety of wavelengths, while Section~\\ref{sec:observations} has details of our observations and data reduction. We analyze the interstellar spectrum of the Cosmic Eye in Section~\\ref{sec:ISM}, and the composite spectrum of its young stellar population in Section~\\ref{sec:stars}. Section~\\ref{sec:intervening} focuses on the numerous intervening absorption systems that potentially confound our interpretation of the galaxy's intrinsic features.  In Section~\\ref{sec:discussion} we discuss the implications of our findings; finally, we summarize our main conclusions in Section~\\ref{sec:conclusions}. Throughout the paper, we adopt a cosmology with $\\Omega_{\\rm M} = 0.3$, $\\Omega_{\\Lambda} = 0.7$, and $H_0 = 70$\\,km~s$^{-1}$~Mpc$^{-1}$. ",
        "conclusions": "\\label{sec:conclusions}  Strong gravitational lensing of a $z = 3.07331$ star-forming [${\\rm SFR} \\simeq 50\\,M_\\odot$~yr$^{-1}$ for a Chabrier (2003) IMF] galaxy magnifies it by a factor of $\\sim 25$ and distorts its image into two $3^{\\prime\\prime}$ long arcs which have been collectively named the `Cosmic Eye' by their discoverers, Smail et al. (2007). The high magnification has allowed us to use the ESI spectrograph on the Keck\\,{\\sc ii} telescope to record the galaxy's rest-frame UV spectrum with high resolution and S/N ratio. By analyzing these data together with existing observations of the Cosmic Eye at other wavelengths, we have reached the following main conclusions.  (i) The interstellar absorption lines exhibit two components, of approximately equal strength, which are respectively blueshifted by  $-70$\\,\\kms\\ and redshifted by $+350$\\,\\kms\\ relative to the stars and H\\,{\\sc ii} regions. While these values apply to the gas with the highest apparent optical depths, both components include absorption spanning  several hundred \\kms. We associate the blueshifted component with a galaxy-wide outflow similar to, but possibly weaker than, those seen in most star-forming galaxies at $z = 2$--3. The redshifted absorption is very unusual, and may represent gas ejected by a previous episode of star formation and now falling back onto the galaxy, or a merger viewed along a favourable line of sight. Alternatively, it may just be a chance superposition of another galaxy along the line of sight.  (ii) Both components of the metal absorption lines show indications that they do not fully cover the OB stars against which they are being viewed; we estimate covering fractions of $\\sim 70$\\% and $\\sim 85$\\%  for the blueshifted and redshifted component respectively. There must also be more pervasive diffuse gas, because the strong damped \\lya\\ line, corresponding to a column density $N$(H\\,{\\sc i})\\,$=  (3.0 \\pm 0.8) \\times 10^{21}$\\,cm$^{-2}$, covers at least 95\\% of the UV stellar continuum. We tentatively associate this high column density of H\\,{\\sc i} with the redshifted component of the metal lines, where absorption from ionized species is weak or missing altogether, and propose that it provides the `foreground screen' of dust responsible for the lower-than-expected far-infrared luminosity of the Cosmic Eye found by Siana et al. (2009) with the \\textit{Spitzer Space Telescope}.  (iii) The internal kinematics of the galaxy lensed into the Cosmic Eye are very complex, our data now adding outflow, and possibly even inflow, to the rotation and velocity dispersion already known from the integral field spectroscopy with adaptive optics by Stark et al. (2008). Ordered rotation, chaotic motions, and outflow all seem to be of comparable magnitude, with $v_{\\rm rot} \\sin i \\approx \\sigma_0 \\approx v_{\\rm blue} \\simeq  50$--70\\,\\kms. However, we do not have a model yet of how these different motions fit together into one coherent kinematic picture.  (iv) Turning to the stellar spectrum, we find that the C\\,{\\sc iv} P~Cygni profile is well fit by a \\textit{Starburst99} stellar population model spectrum having continous star formation with a Salpeter IMF, stellar masses from 5 to 100\\,\\msun, and a LMC/SMC metallicity of $Z \\sim 0.4$\\,\\Zsun. The P~Cygni profiles of the Cosmic Eye, the Cosmic Horseshoe, and MS\\,1512-cB58, three high redshift star-forming galaxies studied at high spectral resolution, are all nearly identical. This is not unexpected, however, when we consider that in each case we see the integrated light of several hundred thousand O-type stars, and that these three galaxies have similar metallicities and dynamical timescales over which star formation is taking place.  (v) The metallicity $Z \\simeq 0.4 Z_\\odot$ deduced for the O stars in the Cosmic Eye is lower by a factor of $\\sim 2$ than the value derived from the analysis of the strong nebular lines from its H\\,{\\sc ii} regions using the R23 index. We consider this apparent discrepancy to reflect systematic offsets between different abundance estimators, rather than intrinsic inhomogeneities in the chemical composition of stars and ionized gas.  (vi) The interpretation of both interstellar and stellar features in the UV spectrum of the Cosmic Eye is complicated by the presence of numerous intervening absorption lines associated with eight absorption systems at redshifts $z_{\\rm abs} = 2.4563$--3.0528. These narrow features are not resolved in existing low-resolution data, highlighting the caution that should be exercised in interpreting the spectra that are typically available for unlensed LBGs.  In closing, the new data presented here further emphasize the complexity of the physical conditions which prevailed in actively star-forming galaxies at redshifts $z = 2$--3. It is remarkable that the three strongly-lensed galaxies targeted by our ESI observations, while showing very similar young stellar populations, are all different in the detailed properties of their interstellar media. Whether such variety is simply the result of different geometries and viewing angles, or has its roots in more fundamental physical reasons remains to be established. Fortunately, with the increasing attention being given to detailed studies of gravitationally lensed galaxies, we can look forward with optimism to a more comprehensive empirical picture of galaxy formation coming together in the years ahead."
    },
    "0910/0910.5462_arXiv.txt": {
        "abstract": "We comment on the calculation mistake in the paper ``$w$ and $w'$ of scalar field models of dark energy'' by Takeshi Chiba, where $w$ is the dark energy equation of state and $w'$ is the time derivative of $w$ in units of the Hubble time. The author made a mistake while rewriting the phantom equation of motion, which led to an incorrect generic bound for the phantom model and an incorrect bound for the tracker phantom model on the $w\\textrm{--}w'$ plane. ",
        "introduction": "\\subsection{Generic bound} The phantom model is scalar field dark energy having negative kinetic term and its equation of state $w<-1$. The energy density and the pressure are given by $\\rho=-\\dot\\phi^2/2+V$ and $p=-\\dot\\phi^2/2-V$, respectively. The equation of motion is given by \\begin{equation}\\label{q1} \\ddot\\phi+3H\\dot\\phi-V_{,\\phi}=0 \\ . \\end{equation} Therefore, the phantom field tends to roll up the potential $V$. The phantom equation of motion can be rewritten as~\\cite{Kujat:2006vj} \\begin{equation}\\label{q2} \\pm {V_{,\\phi}\\over V}=\\sqrt{{-3\\kappa^2(1+w)\\over \\Omega_{\\phi}}} \\left(1+{x'\\over 6}\\right) \\ , \\end{equation} where the plus sign corresponds to $\\dot\\phi>0$ and the minus sign to the opposite, $\\kappa^2=8\\pi G$ and $\\Omega_{\\phi}$ is the fractional energy density of the phantom field. The variable $x$ is defined as \\begin{equation}\\label{q3} x=\\ln\\left(-{1+w\\over 1-w}\\right) \\ , \\end{equation} and $x'$ is the derivative of $x$ with respect to $\\ln a$, which is related with $w'$ as \\begin{equation}\\label{q4} x'={2w'\\over (1-w)(1+w)} \\ . \\end{equation} Since the left-hand side of Eq.~(\\ref{q2}) is positive for the up-rolling phantom field, we have $1+x'/6>0$. Therefore, using Eq.~(\\ref{q4}), the upper bound on $w'$ is obtained as \\begin{equation}\\label{q5} w'<-3(1-w)(1+w). \\end{equation} Note that the upper bound on $w'$ for phantom is smoothly connected to the lower bound on $w'$ for quintessence, the latter of which is shown as Eq.~(2.6) in~\\cite{Chiba:2005tj}.  In~\\cite{Chiba:2005tj}, there is a calculation mistake that the rewritten equation of motion is obtained as \\begin{equation}\\label{q6} \\mp {V_{,\\phi}\\over V}=\\sqrt{{-3\\kappa^2(1+w)\\over \\Omega_{\\phi}}} \\left(-1+{x'\\over 6}\\right) \\ . \\end{equation} As a consequence, an incorrect bound on $w'$ is obtained as \\begin{equation}\\label{q7} w'>3(1-w)(1+w). \\end{equation}  \\subsection{Tracker phantom}\\label{TP} Following the approach in~\\cite{Chiba:2005tj}, the bound can be tightened for the tracker field. Taking the derivative of Eq.~(\\ref{q2}) with respect to $\\phi$, we obtain the tracker equation for the phantom field \\begin{eqnarray}\\label{q8} \\Gamma-1  &=& \\frac{3(w_B-w)(1-\\Omega_{\\phi})}{(1+w)(6+x')}-\\frac{(1-w)x'} { 2(1+w)(6+x')} \\nonumber \\\\  & \\ & -\\frac{2x''}{(1+w)(6+x')^2} \\ , \\end{eqnarray} where $\\Gamma=VV_{,\\phi\\phi}/V_{,\\phi}^2$, $w_B$ is the equation of state of the background matter, and $x''$ is the second derivative of $x$ with respect to $\\ln a$. Therefore, for the tracker solution where $w$ is nearly constant and $\\Omega_{\\phi}$ is initially negligible, $w$ is given by \\begin{equation}\\label{q9} w={w_B-2(\\Gamma -1)\\over 2(\\Gamma -1) +1} \\ . \\end{equation} Thus $\\Gamma<1/2$ is required for tracker phantom, which has $w<-1$.  Following~\\cite{Chiba:2005tj}, we consider a solution in which initially $w$ follows the tracker solution in Eq.~(\\ref{q9}) and then evolves toward $-1$. Therefore, the tracker $w$ in Eq.~(\\ref{q9}) is a lower bound of $w$. In such solution, $x'$ eventually stops decreasing and then increases back to a value near zero. The minimum value of $x'$, $x_m'$, gives an upper bound on $w'$ via Eq.(\\ref{q4}). To find $x_m'$, we put $x''=0$ and $w_B=0$ in Eq.~(\\ref{q8}) and find that \\begin{equation}\\label{q10} x_m'=-6{w(1-\\Omega_{\\phi})+2(1+w)(\\Gamma -1)\\over (1-w)+2(1+w)(\\Gamma -1)} \\ . \\end{equation} Since $x_m'$ is an increasing function of $w$, a lower bound on $x_m'$ is given by that of $w$, for which we take the tracker $w$ in Eq.~(\\ref{q9}) and obtain \\begin{equation}\\label{q11} x_m'>\\frac{6w\\Omega_{\\phi}}{1-2w}>\\frac{6w}{1-2w} \\ . \\end{equation} With Eq.~(\\ref{q4}), we then obtain the upper bound on $w'$ \\begin{equation}\\label{q12} w'<\\frac{3w(1-w)(1+w)}{1-2w} \\ . \\end{equation}  In \\cite{Chiba:2005tj}, there is a mistake that the tracker equation for the phantom field is given as \\begin{eqnarray}\\label{q13} \\Gamma-1  &=& \\frac{3(w_B-w)(1-\\Omega_{\\phi})}{(1+w)(6-x')}-\\frac{(1-w)x'} { 2(1+w)(6-x')} \\nonumber \\\\  & \\ & +\\frac{2x''}{(1+w)(6-x')^2} \\ . \\end{eqnarray} As a consequence, the resulting $x_m'$ is \\begin{eqnarray}\\label{q14} x_m'&=&-6{w(1-\\Omega_{\\phi})+2(1+w)(\\Gamma -1)\\over (1-w)-2(1+w)(\\Gamma -1)}   \\\\&>&6w\\Omega_{\\phi}>6w. \\end{eqnarray} Therefore, an incorrect bound on $w'$ is given as \\begin{equation}\\label{q15} w'<3w(1-w)(1+w). \\end{equation}  The bounds for the phantom model from \\cite{Chiba:2005tj} are shown in Fig.~\\ref{fig1} and the corrected ones we obtain are shown in Fig.~\\ref{fig2}.   \\begin{figure} \\includegraphics[width=8.3cm]{origional} \\caption{\\label{fig1}Bounds on $w'$ as a function of $w$ for the phantom model from \\cite{Chiba:2005tj}. The solid curve is the generic lower bound of the phantom. The dashed curve is the upper bound of the tracker phantom. The allowed region for the tracker phantom is filled with the grey color. } \\end{figure}  \\begin{figure} \\includegraphics[width=8.3cm]{corrected} \\caption{\\label{fig2}Corrected bounds we obtain on $w'$ as a function of $w$ for the phantom model. The solid curve is the generic upper bound of the phantom. The dashed curve is the upper bound of the tracker phantom. The allowed region for the tracker phantom is filled with the grey color.} \\end{figure} ",
        "conclusions": ""
    },
    "0910/0910.2418_arXiv.txt": {
        "abstract": "We apply a reflection-based model to the best available {\\it XMM-Newton} spectra of X-ray bright UltraLuminous X-ray (ULX) sources (NGC~1313 X--1, NGC~1313 X--2, M~81 X--6, Holmberg~IX X--1, NGC~5408 X--1 and Holmberg~II X--1). A spectral drop is apparent in the data of all the sources at energies 6--7\\,keV. The drop is interpreted here in terms of relativistically-blurred ionized reflection from the accretion disk.  A soft-excess is also detected from these sources (as usually found in the spectra of AGN), with emission from O K and Fe L, in the case of NGC~5408 X--1 and Holmberg~II X--1, which can be understood as features arising from reflection of the disk. Remarkably, ionized disk reflection and the associated powerlaw continuum provide a good description of the broad-band spectrum, including the soft-excess. There is no requirement for thermal emission from the inner disk in the description of the spectra. The black holes of these systems must then be highly spinning, with a spin close to the maximum rate of a maximal spinning black hole. The results require the action of strong light bending in these sources. We suggest that they could be strongly accreting black holes in which most of the energy is extracted from the flow magnetically and released above the disc thereby avoiding the conventional Eddington limit. ",
        "introduction": "Ultra-luminous X-ray (ULX) sources are point like, nonnuclear sources observed at other galaxies, with observed luminosities greater than the Eddington luminosity for a $10\\,{\\rm M}_{\\odot}$ stellar mass black hole (BH), with $L_{X}{\\ge}10^{39}$\\,${\\rm erg\\,s^{-1}}$ \\citep{f1}. The true nature of these objects is still open to debate \\citep{mc1}. One fundamental issue is whether the emission is isotropic or beamed along our line-of-sight. A possible scenario for geometrical beaming involves super-Eddington accretion during phases of thermal-timescale mass transfer \\citep{k1}. Alternatively, if the emission is isotropic and the Eddington limit is not violated, ULX must be fuelled by accretion onto Intermediate-Mass BH (IMBH), with masses in the range 100-10\\,000 ${\\rm M}_{\\odot}$ \\citep{cm1}. Currently, there is no agreement regarding the nature of these sources. It is possible that some ULX appear very luminous because of a combination of moderately high mass, mild beaming and mild super-Eddington emission. It is also possible that ULX are an inhomogeneous population, comprised of both a subsample of IMBH and moderately beamed stellar mass black holes \\citep{f2,mc1}.  NGC~5408 X--1, HOLM~II X--1, M 81 X--9 (which is in the companion-dwarf galaxy Holmberg IX, hereafter called HOLM~IX X--1), the ULX in NGC~1313 and M 81~X--6 are ULX located in dwarf (NGC~5408 X--1, HOLM~II X--1 and HOLM~IX X--1) and spiral (NGC~1313 X--1, NGC~1313 X--2 and M 81 X--6) galaxies, respectively. All these ULX peak in X-ray luminosity above $L_{X}=1{\\times}10^{40}$\\,${\\rm erg\\,s^{-1}}$, thus being excellent targets for testing spectral models to the data. They are nearby and located at distances of $D=4.8,3.5,3.7,3.63,4.50$\\,Mpc (\\citealt{karachentsev02,stobbart06,liu08}), for NGC~5408 X--1, HOLM~IX X--1, NGC~1313 X--1, M~81 X--6 and HOLM~II X--1 respectively. NGC~5408 X--1 is among the few ULX for which (10-200\\,mHz) Quasi-Periodic Oscillations (QPO) have been found (\\citealt{stroh1,dewangan06}). NGC~5408~X-1 exhibits X-ray timing and spectral properties analogous to those exhibited by Galactic stellar-mass black hole in the {\\it very high} or {\\it steep power-law} state \\citep{remi1}, but with the characteristic variability timescales (QPO and break frequencies) consistently scaled down \\citep{stroh1}. For NGC~5408 X--1 the inferred characteristic size for the X-ray emitting region is ${\\approx}$ 100 times larger than the typical inner disk radii of Galactic black holes. The highest signal-to-noise spectra of ULX can often be fitted by composite models of a thermal disk together with a hard-powerlaw tail, with a low measured disk temperature of ${\\approx}0.2$\\,keV (\\citealt{miller03,miller04}). If such factors are entirely due to a higher black hole mass, they imply an IMBH with $M{\\approx}100-10\\,000$\\,${\\rm M}_{\\odot}$. Both the morphology and flux of the optical high-excitation nebulae detected around some ULX probably rule out strong beaming as the origin of the X-ray emission (\\citealt{pakull03,kaaret04,soria1,abolmasov07}).  The detection for most ULX of spectral curvature, in the form of a deficit of photons at energies $E{\\ge}2$\\,keV (\\citealt{roberts05,stobbart06,miyawaki09}) has led to the suggestion that most ULX have spectral properties that do not correspond to any of the accretion states known in Black Hole Binaries (BHB), making it unlikely that ULX are powered by sub-Eddington flows onto an IMBH \\citep{roberts07}. The application of Comptonization models to the data (\\citealt{stobbart06,gladstone09} and references therein) results in strikingly high and low values for the coronal opacity (${\\tau}{\\ge}5$) and the electron temperature ($kT_{e}=1-3$\\,keV), difficult to explain for the expected physical conditions in a corona surrounding the black hole. This is very different to the typical values found for BHB during the {\\it low/hard} state, with spectra dominated by Comptonization \\footnote{It has to be noted that similar values, i.e. $kT_{e}=1-3$\\,keV and ${\\tau}{\\ge}5$, have been found in the description of the spectra of BHB during the {\\it steep powerlaw} state (e.g. GRO~J1655--40 and GRS~1915+105; \\citealt{makishima00,kubota04,ueda09}).}, and appears irreconcilable with the IMBH model, which assumes that they operate as simple scaled-up BHB.  Here, we present an alternative interpretation of the spectral shape, based on a physically-justified model commonly used on other accreting black holes. The soft part of the spectrum ($E{\\le}1$\\,keV) -- the soft X-ray excess -- and the high-energy curvature are just aspects of a reflection spectrum expected from accretion (\\citealt{guilbert88,george91}). A major component of reflection, the broad iron K line, has been found in many Seyfert galaxies (\\citealt{tanaka95,nandra08}), accreting stellar-mass black holes (\\citealt{miller07,reis09a}) and even accreting neutron stars (\\citealt{cackett08,reis09b}). Both the soft excess and the relativistic broad iron K line have recently been demonstrated to be part of the same physical process, i.e., the reaction of the disk to irradiation from a high-energy source, in the Seyfert-1 galaxy 1H 0707-495 \\citep{fabian09}. In this paper we investigate whether reflection models can account for the spectra of these ULX and, if so, we put them in the context of the accreting black holes known so far.  For this study, we choose the ULX with the best available data and with the longest exposure time observations (${\\approx}100$\\,ks of exposure time) of the {\\it XMM-Newton} satellite. By using only the highest quality data from the widest band pass, highest sensitivity instruments available we expect to make good statements on the accretion processes in these ULX. In Section \\ref{observ} we describe the observations and data used, in Section \\ref{spec_anal} we report on the results of the application of fits with the reflection model to the spectra and in Section \\ref{discuss} discuss the results obtained. ",
        "conclusions": "\\label{discuss}   Remarkably, a relativistically-blurred, reflection-dominated model gives a very good fit to the sample of ULX with best quality (100\\,ks of time exposure) {\\it XMM-Newton} data.  Very skewed iron line profiles have been found, implying that the emission region is very close to the black-hole and this fact allowed us to determine their spin.  We find that all the ULX in our sample are close to maximally rotating, with the exception of NGC~5408 X--1 and HOLM~II X--1, for which the steep powerlaw did not allow us to properly determine parameters from the Fe line. We have found that, the sources showing a {\\it reflection dominated} spectrum (NGC~1313 X--1, NGC~1313 X--2 and M~81 X--6) have maximally spinning black-holes.  This fact can be understood if light bending is an important effect for these sources. For these sources, the high-energy source (i.e. the jet or corona) is very close to the black hole and the observed direct continuum (i.e. powerlaw) is very low. This would correspond to {\\tt Regime I} of \\citet{miniutti04}, corresponding to a low height of the primary source, where strong light suffered by the primary radiation dramatically reduces the observed powerlaw emission component at infinity. HOLM~IX X--1 shows a {\\it powerlaw dominated} spectrum with a slightly narrower Fe line profile. This would correspond to {\\tt Regime II} of \\citet{miniutti04}, when the height of the primary source is larger, and the observed direct continuum (i.e. powerlaw) stronger. The line profile is narrower than in regime I because the emissivity profile of the disk is flatter, as result of a nearly isotropic illumination. For NGC~5408 X--1 and HOLM~II X--1, we find a clear reflection component but the steepness of the spectra means that the Fe K features are less well-measured.  The spectral solution we have found here for ULX suggests a new explanation for accretion onto spinning black holes. Our model assumes that the dominant source of radiation is a power-law continuum produced a few gravitational radii above a rapidly spinning black hole.  Little thermal radiation is produced by the disc and is undetected by the current data. We assume that the power for this source is extracted magnetically from the disc and transferred to the emission region by magnetic fields. Some of the power may even be extracted magnetically from the spin of the black hole \\citep{blandford77}. Since radiation is only produced above the disc, radiation pressure need not oppose accretion. Indeed it will help squash the disc and maintain the high surface density required for our relatively low ionization parameters and thus observable reflection. (Note that strong light bending occurs in the region being considered which is very close to a black hole, meaning that there is beaming but in the {\\emph opposite} sense to that normally envisaged.)  The relevant radiation for computing the physical Eddington limit in this situation is the thermal disc radiation, not the power-law continuum and reflection associated with it. Since we detect no such thermal radiation, then the situation may be sub-Eddington, even for stellar mass black holes. Whether this solution can work depends on the extent to which magnetic energy extraction can be clean, in the sense of not requiring considerable thermal energy release. We note that the accretion flow is super-Eddington in the conventional interpretation in which the total energy release is considered (especially since some radiation falls straight into the black hole).  If the above solution is appropriate for these objects then they can either be stellar mass black holes with masses of ${\\approx}10\\Msun$ or IMBH of $100$s$\\Msun$. Their spin is then likely to be natal. If they form through stellar collapse and the progenitor star was highly spinning, then it would be likely that the black hole would remain highly spinning. Many studies have argued that a binary companion can spin up a massive star but the magnitude of the spin up is still a matter of debate (e.g., \\citealt{paczynski98,fryer99,brown00}). Accretion would presumably come from a binary companion. An interesting alternative is that they accrete from a debris disc produced by fallback from the (failed) supernova or collapsar \\citep{li03}. This could account for the high metallicity we infer (note that hydrogen-poor accretion would also slightly raise the Eddington limit).  In summary, some ULX may consist of fast spinning black holes accreting rapidly from a metal-rich disc. Accretion power passes magnetically to the primary emission region which lies just outside the accretion flow. Strong reflection occurs, creating a soft excess and very broad iron-K features in the 0.3-10 keV X-ray spectrum. Sensitive hard X-ray spectra in the 10--50~keV band may test this scenario."
    },
    "0910/0910.4714_arXiv.txt": {
        "abstract": "We study the prospects for detecting neutrino masses from the galaxy angular power spectrum in photometric redshift shells of the Dark Energy Survey (DES) over a volume of $\\sim 20\\, h^{-3}$ Gpc$^3$, combined with the Cosmic Microwave Background (CMB) angular fluctuations expected to be measured from the Planck satellite. We find that  for a $\\Lambda$-CDM concordance model with 7 free parameters in addition to a fiducial neutrino mass of  $M_{\\nu} = 0.24$ eV, we recover from DES\\&Planck  the correct value with uncertainty of $ \\pm 0.12$ eV (95 \\% CL), assuming perfect knowledge of the galaxy biasing. If the fiducial total mass is close to zero, then the upper limit is $0.11$ eV(95 \\% CL). This upper limit  from DES\\&Planck is over 3 times tighter than using Planck alone, as DES breaks the parameter degeneracies in a CMB-only analysis. The analysis utlilizes spherical harmonics up to 300, averaged in bin of 10 to mimic the DES sky coverage. The results  are similar if we supplement DES bands (grizY) with the VISTA Hemisphere Survey (VHS) near infrared band (JHK). The result is robust  to uncertainties in non-linear fluctuations and redshift distortions. However, the result is sensitive to the assumed galaxy biasing schemes and it requires accurate prior knowledge of the biasing. To summarize, if the total neutrino mass in nature greater than 0.1eV, we should be able to detect it with DES\\&Planck, a result with great importance to fundamental Physics. ",
        "introduction": "\\renewcommand{\\thefootnote}{\\fnsymbol{footnote}} \\setcounter{footnote}{1} \\footnotetext{E-mail: lahav@star.ucl.ac.uk}   Neutrinos are so far the only dark matter candidates that we actually know exist.  It is now established from  solar, atmospheric, reactor and accelerator neutrino experiments that neutrinos have non-zero mass, but their absolute masses are still unknown.  Cosmology could provide an upper limit on the sum of neutrino masses \\citep[for review see e.g.][]{elg05,lesg06}. The growth of Fourier modes with comoving wavenumber $k > k_{\\rm nr}$ will be suppressed because of neutrino free-streaming, where \\begin{equation} k_{\\rm nr} = 0.026\\left(\\frac{m_{\\nu}}{1\\;{\\rm eV}}\\right)^{1/2} \\Omega_{\\rm m}^{1/2}\\;h\\,{\\rm Mpc}^{-1}, \\label{eq:knr} \\end{equation} for three equal-mass neutrinos, each with mass $m_\\nu$.   The current mass upper limit, obtained using Cosmic Microwave Background (CMB) WMAP5 data, SN Ia and the BAO from 2dFGRS and SDSS, is $M_{\\nu} \\equiv \\sum m_{\\nu} < 0.61$ eV at 95\\% CL \\citep{koma09}.  The challenge now is to bring down reliably the upper limits to the 0.1 eV level or even detect the neutrino mass.  In this way Cosmology could resolve the mass scale of neutrinos. The new generation of deep wide surveys can play a key role in setting a tight upper limit on the neutrino mass, and possibly detect it if the true neutrino mass is sufficiently high. This is due to the order of magnitude increase in volume of the new surveys.   Here we study specifically the ability to set an upper limit on the neutrino mass from the galaxy clustering expected in the photometric redshift survey  Dark Energy Survey (DES), combined with Planck CMB  measurements.  The reason a survey like DES would be effective is its large volume, $\\sim 20\\, h^{-3}$ Gpc$^3$, and large number of galaxies, $\\sim 300$ million.  Crudely,  for a spectroscopic survey, where accurate redshifts are known, the error on the power spectrum scales with the survey effective volume $V_{\\rm eff}$ as \\begin{equation} \\Delta P(k)/P(k) \\propto 1/{\\sqrt {V_{\\rm eff} }}. \\label{eq:dpkveff} \\end{equation} On the other hand, the suppression is proportional in the linear regime to $f_{\\nu} = \\Omega_{\\nu}/\\Omega_m$ \\citep{hu98,kia08}, where \\begin{equation} \\Omega_{\\nu} = \\frac{\\Sigma_i m_i}{93.14 h^2 ~ \\rm eV}. \\label{eq:onumass} \\end{equation} We expect therefore (when all other cosmological parameters are fixed) that the determination of  the upper limit on the neutrino mass would be inversely proportional to ${\\sqrt  {V_{\\rm eff} } }$. From the 2dF Galaxy spectroscopic redshift survey, covering a volume of roughly 0.2 $(h^{-1} Gpc)^3$, the upper limit on the sum of neutrino mass is about 2 eV  at 95 \\% CL \\citep{elg02}. Had DES been a spectroscopic  survey with volume of about 20  $(h^{-1} Gpc)^3$, i.e. 100 times larger,  we would expect an upper limit of 0.2 eV on the sum of neutrino mass. Our detailed calculation below yields an upper limit of 0.1eV for DES\\&Planck. This is tighter than the above back-of-the-envelope calculation, probably as Planck priors are incorporated, and the effective volumes above are only given crudely.  However, DES is photometric redshift survey, so the radial component of distance to galaxies is significantly degraded, resulting in a poorer estimate of the power spectrum  \\citep[e.g.][]{blk05}. Therefore we prefer in this analysis to quantify the galaxy clustering as angular (spherical harmonic) $C_{\\rm \\ell}$ power spectrum derived in photometric redshift shells which are wide enough relative to the photometric redshift errors, and to derive the resulting upper limits more carefully and quantitatively. We defer the  comparison of $P(k)$ and $C_{\\rm \\ell}$ approaches to future studies. The utility of photometric redshifts is now well-established, with many successful techniques being employed  \\citep[e.g.][]{coll04,abd09}. The cosmological  parameter constraints resulting from future photometric redshift imaging surveys have been simulated by  several authors \\citep[e.g.][]{seo03,dol04,zhan06}. Application to data such as the SDSS LRG samples were given by \\cite{blk07} and \\cite{pad07}. These studies mainly emphasized the detection of  baryon acoustic oscillations in the galaxy clustering pattern. Apart from the specific application to DES,  the present paper illustrates more generally  the determination of neutrino mass from photo-z surveys, to our knowledge for the first time.  This paper is organized as follows. In Sections \\ref{sec:galsur} we summarize the DES and VHS surveys and Planck, in Section \\ref{sec:photz} we present the photometric redshifts for DES and DES\\& the Vista Hemisphere Survey (VHS) combined filters. Sections \\ref{sec:clform} give the  formalism for the galaxy angular power spectrum, and the associated joint likelihood with Planck. Section \\ref{sec:result} presents the results for the basic observational and theoretical scenarios, while Section \\ref{sec:analys} provides extensions of the analysis. An overall discussion is given in Section \\ref{sec:conclu}. ",
        "conclusions": "\\label{sec:conclu}  We study the prospects for detecting neutrino masses from the galaxy angular power spectrum in photo-z shells in the Dark  Energy Survey,  combined with CMB fluctuations as will be measured by Planck. Although the core science case for DES is Dark Energy, we see that DES can provide us with other important extra science, such as neutrino mass. Our main conclusions are:  \\begin{itemize}  \\item We forecast for DES\\&Planck a 2-sigma error of total neutrino mass $\\Delta M_{\\nu} \\approx 0.12$ eV.  If the true neutrino mass is very close to zero, then we can obtain an upper limit of 0.11 eV (95\\% CL).  \\item This upper limit  from DES+Planck is over  3 times tighter than using Planck alone, as DES breaks the parameter degeneracies in a CMB-only analysis.  \\item The results are sensitive to the assumed galaxy biasing, and stand if the galaxy bias in known to within 2 per cent. This is feasible given other analyses of the galaxy bias such as the three point correlation function \\citep{ross07}.  \\item The results are robust  to uncertainties in non-linear fluctuations and redshift distortion.  \\item The results  are similar  if we  supplement DES bands (grizY) with the VISTA Hemisphere Survey (VHS) near infrared band (JHK).  \\end{itemize}  DES can also be used to extract information on neutrino mass via other techniques, e.g. weak gravitational lensing, as considered recently for other imaging surveys  \\citep{kit08, ichi09}. We note that the  level of sensitivity for neutrino mass from DES\\& Planck is of much relevance for comparison with the direct measurement of the neutrino mass from laboratory experiments. E.g. the KATRIN tritium beta decay experiment. Furthermore, the DES \\& Planck measurements can be combined with laboratory experiments to derive more accurate neutrino masses \\citep{host07}."
    },
    "0910/0910.3302_arXiv.txt": {
        "abstract": "We present a  hard X-ray spectrum of unprecedented quality of the Galactic supernova remnant W49B obtained with the {\\it Suzaku} satellite. The spectrum exhibits an unusual structure consisting of a saw-edged bump above 8~keV. This bump cannot be explained by any combination of high-temperature plasmas in ionization equilibrium. We firmly conclude that this bump is caused by the strong radiative recombination continuum (RRC) of iron, detected for the first time in a supernova remnant. The electron temperature derived from the bremsstrahlung continuum shape and the slope of the RRC is $\\sim$1.5~keV. On the other hand, the ionization temperature derived from the observed intensity ratios between the RRC and K$\\alpha$ lines of iron is $\\sim$2.7~keV. These results indicate that the plasma is in a highly overionized state. Volume emission measures independently determined from the fluxes of the thermal and RRC components are consistent with each other, suggesting the same origin of these components. ",
        "introduction": "\\label{sec:introduction} W49B (G43.3-0.2) is a Galactic supernova remnant (SNR) with strong X-ray line emissions from highly ionized atoms. It exhibits centrally filled X-rays inside a bright radio shell with a radius of 100 arcsec (Pye et al. 1984). The distance to W49B remains uncertain. It was estimated to lie at a distance of $\\sim$8~kpc (Radhakrishnan et al. 1972; Moffett \\& Reynolds 1994). Using the spectral distribution of HI absorption, however, Brogan \\& Troland (2001) found no clear evidence that W49B is closer to the sun than W49A, which is located at a distance of $\\sim$11.4~kpc (Gwinn et al. 1992). This may indicate a possible association of W49B with the star-forming region W49A. The near-infrared narrowband images indicate a barrel-shaped structure with coaxial rings, which is suggestive of bipolar wind structures surrounding massive stars (Keohane et al. 2007). They also showed an X-ray image from {\\it Chandra}, which has a jet-like structure along the axis of the barrel. They interpreted these findings as evidence that W49B had exploded inside a wind-blown bubble in a dense molecular cloud.  Using {\\it ASCA} data, Hwang et al. (2000) showed that broadband modeling of the remnant's spectrum required two thermal components (0.2~keV and 2~keV) and significant overabundances of Si, S, Ar, Ca, Fe, and Ni. They confirmed that most of the X-ray emitting plasma was nearly in collisional ionization equilibrium (CIE). They also found evidence for Cr and Mn K$\\alpha$ emission.  Kawasaki et al.\\ (2005) claimed the presence of ``overionized\" plasma in W49B through the analysis of {\\it ASCA} 2.75--6.0~keV spectrum. They measured intensity ratios of the H-like K$\\alpha$ (hereafter Ly$\\alpha$) to He-like K$\\alpha$ (He$\\alpha$) lines of Ar and Ca to obtain the ionization temperature ($kT_z$), and found that $kT_z$ ($\\sim$2.5~keV) is significantly higher than the electron temperature ($kT_{\\rm e}$) determined from the bremsstrahlung continuum shape ($\\sim$1.8~keV). Miceli et al.\\ (2006) adopted the same analysis procedure for the {\\it XMM-Newton} spectrum of the central region but found no evidence for the overionized state.  In this letter, we report the firm evidence for overionized plasma in W49B using data from the X-ray Imaging Spectrometers (XIS: Koyama et al.\\  2007) onboard the {\\it Suzaku} satellite (Mitsuda et al.\\ 2007).  \\begin{figure}[t] \\includegraphics[scale=.60]{f1.eps} \\caption{Vignetting-corrected XIS image of W49B in the 1.5--7~keV band shown on a linear intensity scale. Data from the three active XISs are combined. Gray contours indicate every 10\\% intensity level relative to the peak surface brightness. The red circle and orange annulus indicate the source and background regions, respectively. The XIS field of view is shown by the black square. \\label{fig:image}} \\end{figure} ",
        "conclusions": "\\label{sec:discussion} We have found, for the first time, the strong He-RRC of Fe from W49B. Yamaguchi et al.\\ (2009) recently discovered the H-RRC of Si and S from a middle-aged SNR, IC~443, and hence this is the second discovery of a clear RRC from an SNR. We also discovered RRC-accompanied recombination lines, which may provide good diagnostics for the overionized plasma. We review a quantitative verification of our spectral analysis and discuss the implications of the results.  \\subsection{Contribution of the Recombination Lines} We discuss the validity of the best-fit fluxes of the RRC and recombination lines in Table~1. The recombination cross section of the H-like ions into a level of $n$ is approximately given as shown below (e.g., Nakayama \\& Masai \\ 2001). \\begin{equation} \\sigma_{n} \\propto \\frac{1}{n^3} \\bigg(\\frac{3}{2} \\frac{kT_{\\rm e}}{E_{\\rm edge}}+ \\frac{1}{n^2} \\bigg) ^{-1} \\label{eq:RC-line} \\end{equation} We apply this approximation for the He-like ions, but $\\sigma_{1}$ must be reduced by half because one electron is already at the ground state. In principle, the recombination line flux can be estimated by the branching ratio to various levels, but these processes are very complicated. We, therefore, base our discussion only on a simple picture.  We compare the predicted capture and  observed transition rates normalized with the $n$ = 1 value. We can estimate the capture rates using Equation~2 as $\\sigma_2/\\sigma_1$ = 0.62, $\\sigma_3/\\sigma_1$ = 0.25, and $(\\sigma_4+\\sigma_5+...)/\\sigma_1$ = 0.34. On the other hand, the observed flux rates are He$\\alpha_{\\rm rec}$/He-RRC = 0.47 ($\\pm$0.10), He$\\beta_{\\rm rec}$/He-RRC = 0.053 ($\\pm$0.023), and He$\\gamma$-$\\infty_{\\rm rec}$/He-RRC = 0.19 ($\\pm$0.03).\\footnote{Throughout this paper, all the errors in the parentheses are at 90\\% confidence level.} By taking the fractions between the observed and predicted rates, we obtain 76 ($\\pm$16)\\%, 21 ($\\pm$9)\\%, and 56 ($\\pm$9)\\% for the He$\\alpha_{\\rm rec}$, He$\\beta_{\\rm rec}$, and He$\\gamma$-$\\infty_{\\rm rec}$ lines, respectively.  These fractions for the He$\\beta_{\\rm rec}$ and He$\\gamma$-$\\infty_{\\rm rec}$ lines may be conceivable, because one electron at $n$ = 1 suppresses direct transitions from excited levels ($n \\ge$ 3$\\to$1) for the He-like ions. The fraction of He$\\alpha_{\\rm rec}$ is slightly smaller than expected because there should be a significant contribution of the cascade decay electrons from higher levels ($n \\ge$ 3$\\to$2$\\to$1). The real He$\\alpha_{\\rm rec}$ flux may be somewhat larger due to the uncertainty of the response function, but the contribution of the He$\\alpha_{\\rm rec}$ flux relative to the total He$\\alpha$ is  only $\\lesssim$10\\% and does not affect results and following discussion.  \\subsection{Electron and Ionization Temperatures} The electron temperatures determined by the bremsstrahlung continuum shape and the RRC slope are 1.52 (1.50--1.53) keV and 1.43 (1.30--1.59) keV, respectively. These consistent results indicate a common origin of these emissions.  The ionization temperatures directly reflect the ion fractions and hence can be determined as shown below. From the best-fit model in Table~1, the flux ratios of Ly$\\alpha$/He$\\alpha$, He-RRC/He$\\alpha$, and H-RRC/He-RRC are given as 0.016 (0.014--0.018), 0.12 (0.11--0.13), and 0.023 ($\\le$0.10). In Figure~4, we compare these values with the modeled emissivity ratios derived from the radiation code of Masai (1994) for a plasma of $kT_{\\rm e}$ = 1.5~keV. We obtain $kT_z$ = 2.58 (2.46--2.68)~keV, 2.68 (2.63--2.73) keV, and 2.55 ($\\le$3.65) keV, respectively, for the above ratios. We also compare the Ly$\\alpha$/He$\\alpha$ ratio with that derived from the APEC code (Smith et al. 2001), although this code is valid only for a CIE state. We obtain $kT_z$ = 2.46 (2.39--2.54) keV, which is within the margin of error of the above results. The ionization temperatures ($\\sim$2.7~keV) are significantly higher than the electron temperatures ($\\sim$1.5~keV), indicating that the plasma is in a highly overionized state.  The first hint of overionized plasma in W49B was found by Kawasaki et al. (2005). Although they analyzed different elements (Ar and Ca) in different energy bands (2.75--6.0~keV), and derived $kT_z$ from the Ly$\\alpha$/He$\\alpha$ ratio using the CIE plasma code, their results ($kT_z$ $\\sim$2.5~keV and $kT_{\\rm e}$ $\\sim$1.8~keV) are nearly consistent with ours. Our claim is more essential because it is based on clear detection of the recombination structures.  \\subsection{Iron and Nickel Abundances} The abundances listed in Table~1 are valid only for the CIE state and should be modified in the overionized case. Since no plasma code can be applied to the overionized plasma currently, we make possible modifications using an available APEC code.  The He$\\alpha$ intensity is proportional to $Z\\times \\epsilon(kT_{\\rm e}, kT_z)$, where $Z$ and $\\epsilon(kT_{\\rm e}, kT_z)$ are the abundance of the element (solar) and the total emissivity for the He$\\alpha$ for $kT_{\\rm e}$ and $kT_z$ (cm$^3$s$^{-1}$), respectively. The emissivities of the He-, Li-, Be-, and B-like ions for Fe and the He- and Li-like ions for Ni are modified by multiplying the ion-fraction ratio between $kT_z$ = 2.7~keV and 1.5~keV (Mazzotta et al. 1998). The total emissivity is given by adding those in individual ionization states. Multiplying by $\\epsilon$(1.5~keV, 1.5~keV)/$\\epsilon$(1.5~keV, 2.7~keV), we obtain the real abundances in the overionized state as $Z_{\\rm Fe}$$\\sim$4.9~solar and $Z_{\\rm Ni}$$\\sim$5.2~solar. Both the elements are highly over abundant, indicating an ejecta origin of the plasma.  \\subsection{Volume Emission Measure} To check the consistency of the common origin of the bremsstrahlung and RRC emissions, we compare the volume emission measure (VEM). The VEM is given by  $\\int n_e n_{\\rm H} dV/(4\\pi D^2)$, where $n_e$, $n_{\\rm H}$, $V$, and $D$ are the electron and hydrogen densities (cm$^{-3}$), emitting volume (cm$^3$), and distance to the source (cm), respectively.  The VEM of the VAPEC component ($VEM_{\\rm VAPEC}$) is derived from Table~1 as  $VEM_{\\rm VAPEC}$ = 1.61 (1.60 -- 1.62) $\\times 10^{13}$ cm$^{-5}$. On the other hand, the VEM of the RRC plasma ($VEM_{\\rm RRC}$) is calculated from \\begin{equation} F_{\\rm RRC}=\\alpha _1(kT_{\\rm e}) \\times \\frac{n_{\\rm Fe}}{n_{\\rm H}} \\times \\kappa_{\\rm H-like}(kT_z) \\times VEM_{\\rm RRC}, \\label{eq:rrcflux} \\end{equation} where $F_{\\rm RRC}$, $\\alpha _1(kT_{\\rm e})$, $n_{\\rm Fe}$, and $\\kappa_{\\rm H-like}(kT_z)$ are the flux of the He-RRC (cm$^{-2}$s$^{-1}$), the K-shell recombination rate coefficient for $kT_{\\rm e}$ (cm$^{3}$s$^{-1}$), the number density of Fe (cm$^{-3}$), and the ion fraction of H-like Fe for $kT_z$, respectively. According to Badnell (2006), the total radiative recombination rate of He-like Fe at $kT_{\\rm e}$ = 1.5~keV is given as $\\sim$3.9$\\times 10^{-12}$ cm$^{3}$s$^{-1}$. The recombination rate into the ground state is given using Equation~2 as $\\sigma_{1}/(\\sigma_1 + \\sigma_2 +...)\\sim$0.45. The value of $n_{\\rm Fe}/n_{\\rm H}$ is calculated using $Z_{\\rm Fe}$ in Section 4.3 and the number density of the solar photosphere (Anders \\& Grevesse 1989). Using the observed He-RRC flux, we obtain\\footnote{Here, we consider that the $kT_z$ error is 2.4--2.8~keV. The large error of $VEM_{\\rm RRC}$ is due to this effect.} $VEM_{\\rm RRC}$ = 0.81 (0.58 -- 1.60)$\\times 10^{13}$ cm$^{-5}$. The two independent estimations of VEM give consistent results, supporting the same origin of the overionized plasma in W49B.  The origin of the overionized plasma in SNRs is an open question, and beyond the scope of this paper. We simply note the possibility of the cooling of electrons via thermal conduction (Kawasaki et al. 2005) or a more drastic cooling caused when the blast wave breaks out of some ambient matter into the rarefied interstellar medium (Yamaguchi et al. 2009). If the latter is the case, a massive progenitor that had blown a stellar wind is likely to be favored.  \\begin{figure}[t] \\includegraphics[scale=.60]{f4.eps} \\caption{Predicted emissivity ratios of Ly$\\alpha$/He$\\alpha$ (black), He-RRC/He$\\alpha$ (red), and H-RRC/He-RRC (green) of Fe as a function of the ionization temperature ($kT_z$) for an electron temperature of 1.5~keV (Masai 1994). The horizontal dashed lines represent 90\\% errors of the observed values. \\label{fig:ratio}} \\end{figure}"
    },
    "0910/0910.1594_arXiv.txt": {
        "abstract": "{% The formation of early-type dwarf (dE) galaxies, the most numerous objects in clusters, is believed to be closely connected to the physical processes that drive galaxy cluster evolution, like galaxy harassment and ram-pressure stripping. However, the actual significance of each mechanism for building the observed cluster dE population is yet unknown. Several distinct dE subclasses were identified, which show significant differences in their shape, stellar content, and distribution within the cluster. Does this diversity imply that dEs originate from various formation channels? Does ``cosmological'' formation play a role as well? I try to touch on these questions in this brief overview of dEs in galaxy clusters. } ",
        "introduction": "\\begin{figure} \\centering \\includegraphics[width=50mm]{dEbigsmall_2_s.eps} \\caption{Two Virgo cluster dEs (VCC\\,856 \\& 839) in a deep exposure with ESO 2.2m/WFI. The image shows an area of $6'\\times5'$. } \\label{dEbigsmallfig} \\end{figure}  Early-type dwarf (dE) galaxies play a key role in understanding galaxy cluster evolution. Their low mass and low density make them more susceptible to physical effects than giant galaxies, and they are a large population, outnumbering all other galaxy types in dense environments by far. They are thus ideal probes of the mechanisms that govern galaxy formation and evolution in a cluster environment. Nevertheless, dEs are also interesting in their own right. Despite their rather unspectacular appearance, a surprising complexity in their characteristics has become evident, in terms of kinematics (rotating vs.\\ pressure supported), structure (flat vs.\\ round), and spatial distribution (cluster center vs.\\ outskirts). This zoo of dEs and their possible origin(s) is an ongoing challenge for observers and theorists. ",
        "conclusions": "\\subsection{Early versus late formation}  Most dEs in cluster regions of high density are pressure supported (Fig.~\\ref{rotfig}), whereas in intermediate and low-density regions, more galaxies seem to be flattened by rotation or be rotationally supported\\linebreak \\citep{vanZee2004a,Toloba2009}. Does this, by itself, tell us something about the formation history of those dEs? Simulations show that relatively flat low-mass cluster galaxies experiencing repeated strong tidal interactions become dynamically hotter: their $v/\\sigma$ decreases and they become ``rounder'', i.e.\\ their axial ratio increases\\linebreak \\citep{Mastropietro2005}. Such interactions are more frequent in the central region of a cluster. Thus, if all dEs once had similar levels of rotational support, those in the central cluster region would have become dynamically hotter, in agreement with the observational findings (I call this {\\it case 1}). Likewise, if those that are now closer to the center have resided in the cluster for a longer time, they experienced more such interactions, leading to the same result ({\\it case 2}). Similar scenarios apply to the observation that galaxies in the central cluster regions have higher stellar population ages on average \\citep{Smith2008}: ram-pressure stripping, and also tidal stripping, are stronger in the central region, thus quenching star formation earlier in the galaxies that reside there. But without further considerations, it cannot be decided whether those galaxies have spent a longer time in the cluster, and were therefore stripped earlier (case 2), or whether they were simply stripped more efficiently (case 1).  From case 2, one might be tempted to conclude a later infall and transformation of the progenitors of those dEs that are now in the cluster outskirts and are significantly flattened by rotation. However, case 1 does not imply anything about when the dEs, or their progenitors, have entered the cluster, or whether they have always been part of the cluster since its formation phase. This is the point at which input from the theoretical side is essential: for case 2, we need to have an understanding of the timescale between the infall of a group of progenitor galaxies and the point when these galaxies (or their remnants) reach a stage in which they are concentrated towards the cluster center and show a peaked and fairly regular distribution of heliocentric velocities. For case 1, probably even more complex input physics is needed, since semi-analytical model predictions for ``cosmological dwarfs'' that form in dense environments need to be tested against the observations. Equally important are observational studies of dEs at significantly earlier epochs, and comparisons to the local population (cf.\\ \\citealt{Andreon2008}; \\citealt{Harsono2009}, also this issue).  \\subsection{Multiple origins?}  Would it be possible, despite the complex characteristics and several different subclasses, to explain all dEs by just a single formation channel? One could, for example, imagine that dEs with disks (dE(di)s) do not form an intrinsically different subclass, but instead they simply constitute the flat tail of the other dEs, with their disk components having not been destroyed yet. The fraction of dE(di)s with nuclei ($\\sim$75\\%, \\citealt{p4}) agrees with  the overall fraction of nucleated dEs in the (bright) magnitude range of the dE(di)s. However, while bright non-nucleated dEs (without disks or blue centers) are significantly flatter than faint ones, this is not the case for the nucleated dEs (again without disks or blue centers). Given the lower stellar population age of dE(nN)s \\citep{Rakos2004}, and their location in regions of lower density, a possible explanation would be that the brighter dE(nN)s still need to experience further dynamical heating through tidal interactions, before they become as round as the dE(N)s. For the faint dE(nN)s, though, one would have to assume that they were more susceptible to such interactions, and therefore already have a rounder shape. On the other hand, the correlation of the projected shape of dE(N)s and their heliocentric velocity \\citep{Lisker2009velo} includes both bright and faint objects, and seems to imply that the roundest shapes are achieved only when the galaxies' orbits are significantly circularized.  To pursue the ``unification'' of dE subclasses, another requirement is that the dE(nN)s would need to form nuclei rather soon, before their stellar population ages are comparable to today's dE(N)s. (Or, alternatively, there might have been some mechanism in the past that supported nucleus formation, but is not as efficient anymore today.) Perhaps nuclei are currently being formed in the centers of the dEs with blue cores, where we are witnessing ongoing star formation \\citep{p2}. This would lead to nuclei with younger stellar populations than those of their host galaxies. However, \\citet{Cote2006} found dE nuclei to have intermediate to old ages, and \\citet{Lotz2004} measured similar colours of nuclei and dE globular clusters (GCs). Alternatively, most nuclei could form through coalescence of GCs, which would indeed be a more efficient process in the central cluster regions \\citep{Oh2000}. Still, some effect would have to explain why the ratio of dE(N)s and dE(nN)s increases strongly with increasing dE luminosity \\citep{san85b}, especially if the dE(nN)s were to be the immediate progenitors of the dE(N)s.  Similar to the above argument about nucleus formation, the spatial distribution of dE(nN)s would soon have to become significantly more centrally concentrated.\\linebreak \\citet{Conselice2001} derived a two-body relaxation time for Virgo dEs of much more than a Hubble time. However, two-body relaxation might be too simple for the real situation. These issues are probably best investigated with cosmological simulations, by ``flagging'' different groups or individual galaxies when they enter a cluster, and following the evolution of their combined (observable) state over various timescales. Current and future generations of simulations\\linebreak \\citep[e.g.][]{BoylanKolchin2009} certainly provide such a possibility.  Obviously, many ifs and thens are necessary to explain all dEs by a single formation channel, and several authors have argued in favour of multiple channels (\\citealt{p4}; \\citealt{Poggianti2001}; \\citealt{Rakos2004};\\linebreak \\citealt{vanZee2004b}). Nevertheless, several aspects of dEs still need to be explored both theoretically and observationally, like the presence and characteristics of their globular clusters and what they can tell us about their evolutionary history \\citep[see the discussion in][]{Boselli2008}.  \\subsection{A plea for the study of early-type dwarfs}  Galaxy mass is one of the main drivers of galaxy evolution and appearance. Internal physical processes like supernova feedback \\citep{DekelSilk1986} or the dynamical response to gas loss \\citep{YoshiiArimoto1987} have a stronger impact on dwarf galaxies than on giants, making dwarfs valuable test objects for galaxy formation models\\linebreak \\citep{Janz2008,Janz2009a}. Nevertheless, if we want to use them as probes of the physical mechanisms that govern galaxy evolution, we need to understand the origin(s) of this  most abundant galaxy population of clusters --- which is a difficult task. The intriguing complexity of dEs is contrasted with the moderate amount of high-quality observational data. Now that the largest cosmological simulations and their semi-analytic models begin to reach significantly into the dwarf regime, it is essential to build complete observational samples providing a thorough characterization of fundamental scaling relations at low masses, involving structure, kinematics, and stellar population properties. These will provide indispensable benchmarks for new generations of galaxy models, as well as for future studies of dEs at significantly earlier epochs with extremely large telescopes."
    },
    "0910/0910.4464_arXiv.txt": {
        "abstract": "{Cyclotron resonant scattering features are an essential tool for studying the magnetic field of neutron stars. The fundamental line provides a measure of the field strength, while the harmonic lines provide information about the structure and configuration of the magnetic field. Until now only a handful of sources are known to display more than one cyclotron line and only two of them have shown a series of harmonics.} {The aim of this work is to see the first harmonic cyclotron line, confirming the fundamental line at $\\sim$22 keV, thus increasing the number of sources with detected harmonic cyclotron lines.} {To investigate the presence of absorption or emission lines in the spectra, we have combined \\emph{RXTE} and \\emph{INTEGRAL} spectra. We modeled the 3--100 keV continuum emission with a power law with an exponential cut off and look for the second absorption feature.} {We show evidence of an unknown cyclotron line at $\\sim$47 keV (the first harmonic) in the phase-averaged X-ray spectra of 4U~1538$-$52. This line is detected by several telescopes at different epochs, even though the S/N of each individual spectrum is low.} {We conclude that the line-like absorption is a real feature, and the most straightforward interpretation is that it is the first harmonic, thus making 4U 1538$-$52 the fifth X-ray pulsar with more than one cyclotron line.} ",
        "introduction": "Cyclotron resonant scattering features (CRSFs), usually referred to as `cyclotron lines', have proved to be powerful tools for directly studying the magnetic field in neutron stars. CRSFs are present in the hard X-ray spectra of several X-ray pulsars and originate in the ``cyclotron process'' under extreme conditions.  Through $E_{cyc}=11.6B_{12} \\times (1+z)^{-1}$ keV (the `12-B-12 law', where $z$ is the gravitational redshift), an energy of the fundamental feature in the hard X-rays indicates that the magnetic fields are rather strong ($ B \\sim 10^{12}$\\,G). Under such conditions, the interaction of the electrons and radiation must be treated quantum-mechanically. The behaviour of an electron in a strong magnetic field implies that the electron energy must be quantized in so-called Landau levels. These absorption features stem from the resonant scattering of photons by electrons, also referred to as cyclotron lines.  While the fundamental energy of the cyclotron line provides valuable information about the magnitude of the field, it is only through the detection and the analysis of the harmonic lines that we can get direct information about the geometrical configuration of the B field (\\cite{harding} 1991; \\cite{araya} 2000; \\cite{schonherr07} 2007). However, to date, only in a handful of systems have harmonic lines been discovered, and only two systems have shown more than two (\\cite{santangelo99} 1999 \\& \\cite{coburn05} 2005). It is therefore paramount to add as many systems to this selected group as we can.  In this work, we present a spectral analysis of the high mass X-ray binary pulsar 4U 1538$-$52. It is an eclipsing system consisting of the B0 I supergiant star QV Nor and a neutron star with an orbital period of $\\sim$3.728 days (Clark \\cite{clark}). The orbital eccentricity is $\\sim$0.08 (Corbet et al. \\cite{corbet}), although more recently a higher value of $\\sim$0.17 was deduced by Clark (\\cite{clark}). The X-ray eclipse lasts $\\sim$0.6 days (Becker et al. \\cite{becker}). The system is fairly bright in X-rays. The estimated flux is $\\sim (5-20)\\times 10^{-10}$ erg s$^{-1}$ cm$^{-2}$ in the 3$-$100 keV range (Rodes~\\cite{rodesPhD}). Thus, assuming a distance of the source of $\\sim$5.5 kpc (Becker et al. \\cite{becker}, Parkes et al. \\cite{parkes}) and an isotropic emission, the luminosity follows $\\sim (2-7)\\times 10^{36}$ erg/s. The magnetized neutron star has a spin period of $\\sim$529 s (Davison \\cite{davison}; Becker et al. \\cite{becker}).  \\begin{figure*}[h!t] \\centering \\includegraphics[angle=-90,width=9cm]{12815fg1.ps} \\includegraphics[angle=-90,width=9cm]{12815fg2.ps} \\caption{Combined spectrum and model of data obtained with \\emph{PCA} (3$-$20 keV) and \\emph{HEXTE} (17$-$100 keV). Both data sets belong to the run carried out in 2001 and their orbital phases are 0.53 and 0.66, respectively. Bottom panels show the residuals in units of $\\sigma$ with respect to the model (see Section~\\ref{RXTEanalysis} for details). } \\label{rxte} \\end{figure*}  The pulse-phase averaged X-ray spectrum of 4U 1538$-$52 has usually been described either by an absorbed power law modified by a high energy cutoff, a power law modified by a Fermi-Dirac cutoff, or by two power laws with indices of opposite sign multiplied by an exponential cutoff (the NPEX model, Mihara~\\cite{mihara}; Rodes et al. \\cite{rodes1}). In addition to these continuum models, an iron fluorescence line at $\\sim$6.4 keV and a cyclotron resonant scattering feature at $\\sim$20 keV discovered by \\emph{Ginga} (Clark et al. \\cite{clark90}) are needed to describe the data. The variability of this CRSF was studied by Rodes-Roca et al. (\\cite{rodes2}). Rossi X-ray Timing Explorer (\\emph{RXTE}) (Coburn \\cite{coburn}) and \\emph{BeppoSAX} data (Robba et al.  \\cite{robba}) did not show evidence of the first harmonic at $\\sim$40 keV. Robba et al. (\\cite{robba}) presented some evidence of an absorption feature around 50 keV; however, because of the lack of a signal-to-noise ratio of the spectrum at these energies, the feature could not be confirmed.  In this paper, we report on the 3--100 keV analysis based on the observations of 4U 1538$-$52 performed by the \\emph{RXTE} and \\emph{INTEGRAL} satellites. In Sect.~\\ref{data} we describe the observations and data analysis. In Sect.~\\ref{analyse} spectral analysis are presented, and summarized in Sect.~\\ref{conclusion}. ",
        "conclusions": "\\label{conclusion}  We presented the spectral analysis of 4U 1538$-$52 using data from \\emph{RXTE} and \\emph{INTEGRAL}. We present evidence for a previously unknown absorption line like feature in the phase-averaged spectrum of the source.  As we have shown in Section~\\ref{analyse}, we have been able to achieve a good fit to the phase-averaged spectra by including a Lorentzian absorption line at $\\sim$47 keV into the model (see Fig.~\\ref{ftest}). This absorption line is clearly visible whenever the signal-to-noise ratio in the spectrum is good enough to allow an analysis of the data. The most straightforward interpretation for this feature is that it is the first harmonic of the $\\sim$ 22\\,keV fundamental CRSF.  According to the theory, cyclotron lines are due to the resonant scattering of photons by electrons whose energies are quantized into Landau levels by the magnetic field (M\\'esz\\'aros \\cite{meszaros}). The quantized energy levels of the electrons are harmonically spaced in the first order, such that the first harmonic line should be placed at twice the energy of the fundamental line, i.e. $ 2\\times E_{\\rm cyc}\\approx 42$ keV. In reality, however, the coupling factor between the fundamental and first harmonic is with $\\sim 2.20$ slightly higher than 2.0.  This anharmonic spacing, however, has been observed already in several systems where more than one line is present. As explained by \\cite{schonherr07} (2007), the relativistic photon-electron scattering already produces some anharmonicity, because photons with energies close to the Landau levels may not escape the plasma if their energies are not changed by inelastic scattering. This, however, can not be the only reason because some systems show an anharmonic spacing larger than that predicted by this effect. A possibility to explain this difference is to take into account that the optical depths of the fundamental and the first harmonic could be different if they are formed at different heights above the neutron star. With increasing height, the strength of the magnetic field decreases resulting in a different CRSF energy. Another possibility is to consider a displacement of the magnetic dipole which would also explain the difference of energy of the two lines if the lines originate from the different poles of the neutron star. Therefore a significant phase dependence of the strength of the both lines is expected, however, the low signal-to-noise ratio at higher energies prevents us to test this hypothesis with the current data sets."
    },
    "0910/0910.1846_arXiv.txt": {
        "abstract": "We present new advances in the spectral extraction of point-like sources adapted to the \\textit{Infrared Spectrograph} onboard the \\textit{Spitzer Space Telescope}. For the first time, we created a super-sampled point spread function of the low-resolution modules. We describe how to use the point spread function to perform optimal extraction of a single source and of multiple sources within the slit. We also examine the case of the optimal extraction of one or several sources with a complex background. The new algorithms are gathered in a plugin called \\texttt{AdOpt} which is part of the \\texttt{SMART} data analysis software. ",
        "introduction": "The ideal spectral extraction algorithm for point-like sources yields the maximum signal-to-noise ratio (S/N) while at the same time preserving the spectrophotometric fidelity. The standard extraction method is based on co-adding the flux in the cross-dispersion direction within a window large enough to contain (most of) the source flux. This method is known as ``tapered-column\" in the  ``Spectroscopy Modeling Analysis and Reduction Tool\" (\\texttt{SMART}\\footnote{\\texttt{SMART} is available at \\textit{http://isc.astro.cornell.edu/smart/}.}, Higdon et al.\\ 2004) and is equivalent to the ``regular extraction\" of the \\textit{Spitzer} IRS Custom Extractor (\\texttt{SPICE}), provided by the \\textit{Spitzer Science Center} (\\textit{SSC}). Although this generally produces satisfactory results, the inclusion of noisy pixels which do not contain a significant fraction of the flux inevitably tends to degrade the quality of the extracted spectrum. This is because every pixel in the extraction window is given the same weight. Moreover, the extraction window is part of a pseudo-rectangle which is defined as a zone in the detector array where the wavelength is uniform (Figure\\,\\ref{fig:pseudo}). When a quadrilateral boundary crosses a pixel, the signal is assumed to be evenly distributed within that pixel, which can lead to artificial flux variations in the extracted spectrum. Because of the interplay of the angled spectral trace and the widening extraction aperture, this error oscillates with wavelength with an amplitude which decreases toward longer wavelengths.   Knowledge of the point spread function (PSF) of the instrument can solve these problems, as it enables the so-called optimal extraction technique which weights the extracted data by the S/N of each pixel (Horne 1986). Optimal extraction therefore significantly reduces the statistical noise in the final spectrum compared to more typical extraction algorithms, which weight all the data within the window equally. So far, efforts on optimal extraction for the \\textit{IRS} have made use of template PSFs or analytical PSFs to fit the cross-dispersion profile of the data: \\begin{itemize} \\item Virtually any  spectrophotometric standard star can be used to derive a template PSF, which then allows an empirical estimate of the PSF at the reference positions (referred to as ``\\textit{nod}\" positions\\footnote{See the \\textit{Spitzer/IRS} observer's manual at \\textit{http://ssc.spitzer.caltech.edu/documents/SOM/}}). This method is currently used by the software \\texttt{SPICE} (see Narron et al.\\ 2007). The resulting extraction undeniably provides better S/N as compared to a tapered column extraction, although the corresponding algorithm is sensitive to cross-dispersion offsets between a source position and the nod position. A simple shift of the PSF is not sufficient to acquire the best S/N possible, and in some cases, low-frequency oscillations can appear in the spectrum due to the misalignment. The reasons are inherent to the data used to compute the PSF template since the latter is created in a specific observational mode (default positions along the slit, default data sampling). \\item Analytical PSFs allow estimating the instrumental profile for any position along the slits. The independent effort from the ``Core to Disks\" legacy program (Evans et al.\\ 2003) uses such PSFs along with an extended emission background which is determined on-the-fly (see Lahuis et al.\\ 2007). Analytical PSFs can bear some uncertainties due to the lack of knowledge of the exact instrumental profile. \\end{itemize}  \\begin{figure}[b!] \\epsscale{1.0} \\plotone{fig1.eps} \\caption{The standard extraction aperture for a single wavelength, superimposed on the grid defined by the pixels on the detector array. The quadrilateral represents the extraction window, which is a ``tapered-column\" whose width increases proportionally with wavelength to account for the varying point spread function. The tilt of the quadrilateral reflects the fact that the row axis of the detector array is not parallel to a line of constant wavelength.\\label{fig:pseudo}} \\end{figure}  We have developed a new optimal extraction algorithm to be used with either of the low-resolution modules of the \\textit{Infrared Spectrograph} (\\textit{IRS}, Houck et al.\\ 2004) onboard the \\textit{Spitzer Space Telescope} (Werner et al.\\ 2004). The new extraction algorithm is available \\textit{via} the plugin \\texttt{AdOpt}\\footnote{The documentation is available at \\textit{http://isc.astro.cornell.edu/SmartDoc/SmartOptimal}.}, part of the new release of \\texttt{SMART}. In a nutshell, an empirical super-sampled PSF has been constructed for each row of the detector array and a multi-linear regression algorithm is used to weight the pixels and derive the flux from the source. The optimal extraction is based on detector rows rather than pseudo-rectangles to treat each pixel as indivisible and thus avoid uncertainties due to the lack of knowledge on the pixel response function. The algorithm presented in this paper produces a considerably higher S/N, up to a factor of $\\sim2$ for faint sources, than the current extraction method available in \\texttt{SMART} for point-like sources (``tapered column\"). An additional advantage of using a super-sampled PSF is that it remains valid anywhere along the aperture in the cross-dispersion direction. We found that the optimal extraction method is extremely sensitive to offsets between a source position and the nominal position. For this reason a new algorithm to locate the source in the slit has been implemented, with a precision of better than a twentieth of a pixel. It is also now possible to extract spectra of spatially blended sources, which is crucial when dealing with crowded regions, such as stellar clusters or nearby galaxies. Such as what can be achieved using iterative techniques (Lucy \\& Walsh 2003), the cross-dispersion profile of the data is decomposed into its components, including the spatial profiles of any number of sources in the slit and the extended background emission. Finally, it is also possible to extract sources with significant offset in the dispersion direction by calculating a modified PSF on-the-fly.   In the following section the basic steps to construct the PSF are given. Section\\,\\ref{sec:method} details the mathematical description of the method and its application for extracting multiple sources are detailed. Section\\,\\ref{sec:applications} briefly explains some specific applications. ",
        "conclusions": "We have presented the optimal extraction algorithm adapted to the \\textit{Infrared Spectrograph} onboard \\textit{Spitzer}. Besides providing a significant increase in the signal-to-noise ratio as compared to the regular ``tapered\" extraction, the new algorithm allows extraction of multiple sources at any location within the aperture. The optimal extraction is also implemented for sources shifted in the dispersion direction. Finally, it is possible to perform an optimal extraction with a complex extended background emission.  The \\texttt{AdOpt} plugin is released with the \\texttt{SMART} package (versions equal or later than 8.0). The code is available at the Infrared Science Center website, along with extensive documentation. \\texttt{SMART} and the optimal extraction will be maintained in the future, with a possible inclusion of the optimal extraction for the high-resolution modules."
    },
    "0910/0910.1241_arXiv.txt": {
        "abstract": "Using new and published photometric observations of $\\mu^{1}~{\\mbox{Sco}}$~(HR~6247), spanning 70~years, a period of 1.4462700(5)~days was determined. It was found that the epoch of primary minimum suggested by Shobbrook at HJD~2449534.178 requires an adjustment to HJD~2449534.17700(9) to align all the available photometric datasets. Using the resulting combined-data light-curve, radial velocities derived from IUE data and the modelling software \\phoebe{}, a new system solution for this binary was obtained. It appears that the secondary is close to, or just filling, its Roche-lobe. ",
        "introduction": "$\\mu^{1}~{\\mbox{Sco}}$ (HR~6247; HD~151890; HIP~82514) was only the third spectroscopic binary to be discovered and is listed in the GCVS \\citep{Samus04} as an eclipsing binary variable. Over the intervening years there have been a number of detailed measurements and several significant studies. It is now classified as a semi-detached~(sd) binary as one component is believed to fill its Roche lobe. \\citet{CesterFedel77} considered $\\mu^{1}~{\\mbox{Sco}}$, to be unusual owing to: \\begin{enumerate} \\item their determination of an apparently high mass-ratio; \\item indications that the secondary component has a larger radius than the primary and overflows its Roche lobe; \\item both components appear to lie on the main sequence; \\item it being a member of the Scorpius-Centaurus cluster thus indicating an age of no more than $10^7$~years hence the possibility this system has just arrived on the main sequence. \\end{enumerate}  A comprehensive analysis of $\\mu^{1}~{\\mbox{Sco}}$ was undertaken to determine the best estimate of period from all available photometric data, establish a suitable epoch and calculate the key parameters defining the system.  The photometric data used included published measurements from: \\begin{enumerate} \\item \\citet{RudnickElvey38}, \\item \\citet{vanGent39}, \\item \\citet{Stibbs48}, \\item Tycho collected between~1990 and~1993 \\citep{ESA97,OchsenbeinBauer00}, \\item Hipparcos collected between~1990 and~1993 \\citep{ESA97,OchsenbeinBauer00}, \\item \\citet{Shobbrook04}, \\end{enumerate} and new measurements taken in 2006, 2007 and 2008 by one of the authors~(Moon).  The photometric data was then combined with radial velocities obtained by \\citet{SticklandSahade96} from IUE spectral data~\\citep{IUE09} and the resulting dataset analysed using the software program \\phoebe{} \\citep{Prsa03,PrsaZwitter05c,PrsaGuinan08,Phoebe09} based on the established \\citet{WilsonDevinney71} theoretical construct \\citep{KallrathMilone99}. ",
        "conclusions": "The light variations of $\\mu^{1}~{\\mbox{Sco}}$ appear to be stable over the 70~years for which photometric data exists. Using the comprehensive dataset of more than 500~photoelectric measurements assembled here, the period was determined to be 1.4462700(5)~day. A suitable epoch based on the primary minimum is suggested to be HJD~2449534.17700(9).  Using all readily available data, and the \\phoebe{} software package, the data for $\\mu^{1}~{\\mbox{Sco}}$ could be best represented by considering the system to be a semi-detached binary in which the size of the secondary is close to, or fills, its Roche Lobe. The system properties determined from comprehensive, iterative modelling using \\phoebe{} indicate that the masses of the components may be somewhat less than previously thought. With the secondary appearing to fill its Roche lobe it maybe possible to obtain spectroscopic evidence of gaseous streams in the $\\mu^{1}~{\\mbox{Sco}}$ system noting that Doppler tomography has produced indirect images of gas flows in interacting binaries \\citep{Richards06}. However, for Algol-type binaries the gaseous streams are faint relative to the main sequence primary star and their detection challenges current techniques. Improvements to techniques would likely be required to detect such gas flows in $\\mu^{1}~{\\mbox{Sco}}$.  There are currently only limited radial velocity data available for the $\\mu^{1}~{\\mbox{Sco}}$. Given that the radial velocity data provides initial values for parameters such as the semi-major axis, masses and radii, further data could be useful in improving the modelling of this binary system. Additionally, photometric measurements near phases 0.1, 0.4, 0.6 and 0.9 could aid refinement of the current theoretical model through better defining some apparent features.  With the semi-major axis of the orbit subtending 0.39~mas, this star may be a suitable target for the Sydney University Stellar Interferometer~(SUSI)."
    },
    "0910/0910.3250_arXiv.txt": {
        "abstract": "An accurate geometric distortion solution for the \\textit{Hubble Space Telescope} UVIS-channel of Wide Field Camera 3 is the first step towards its use for high precision astrometry. In this work we present an average correction that enables a relative astrometric accuracy of $\\sim$1 mas (in each axis for well exposed stars) in three broad-band ultraviolet filters (F225W, F275W, and F336W).  More data and a better understanding of the instrument are required to constrain the solution to a higher level of accuracy. ",
        "introduction": "\\label{sec1}  The accuracy and the stability of the geometric distortion\\footnote{ A specification is needed. With the term ``geometric distortion'', or GD, which we will use hereafter, we are lumping together several effects under the same term: the optical field-angle distortion introduced by camera optics, light-path deviations caused by the filters, non-flat CCDs, alignment errors of CCDs on the focal plane, etc..}  (GD) correction of an instrument is at the basis of its use for high precision astrometry.  The particularly advantageous conditions of the {\\it Hubble Space Telescope (HST)} observatory make it ideal for imaging-astrometry of (faint) point sources.  The point-spread functions (PSFs) are not only sharp and (essentially) close to the diffraction limit --which directly results in high precision positioning-- but also the observations are not plagued by atmospheric effects (such as differential refraction, image motion, differential chromatic refraction, etc), which severely limit ground-based astrometry.  In addition to this, {\\it HST} observations do not suffer from gravity-induced flexures on the structures of the telescope (and camera), which add (relatively) large instabilities in the GD of ground-based images, and make its corrections more uncertain.  Last May 14, the brand-new {\\it Wide Field Camera 3} (WFC3) was successfully installed during the {\\it Hubble Servicing Mission 4} (SM4, May 12-24 2009).  After a period of intense testing, fine-tuning, and basic calibration, last September 9th, 2009, the first calibration- and science-demonstration images were finally made public.  Our group is active in bringing {\\it HST} to the {\\it state of the art} of its astrometric capabilities, that we used for a number of scientific applications (e.g.\\ from King et al.\\ 1998, to Bedin et al.\\ 2009, first and last accepted papers).  Now that the {\\it ``old''} ACS/WFC is successfully repaired, and that the new instruments are installed, our first step is to extend our astrometric tools to the new instruments (and to monitor the old ones).  This paper is focused on the geometric distortion correction of the \\textit{UV/Optical} \\textit{(UVIS)} channel of the WFC3.  Since the results of these efforts might have some immediate public utility (e.g.\\ relative astrometry in general, stacking of images, UV-identification of X-counterparts such as pulsars and CVs in globular clusters, etc.), we made our results available to the WFC3/UVIS user-community.  We immediately focused our attention on a deep UV-survey of the core of the Galactic globular cluster $\\omega$~Centauri (NGC~5139), where some well dithered images were collected.  The dense --and relatively flat-- stellar field makes the calibration particularly suitable for deriving and monitoring the GD on a relatively small spatial scale. In addition, while most of the efforts to derive a GD correction will be concentrated on relatively redder filters, we undertook a study to determine the GD solutions of the three bluest broad-band filters (with the exception of F218W): F225W, F275, and F336W.  The WFC3/UVIS layout is almost indistinguishable from that of ACS/WFC\\footnote{\\sf http://www.stsci.edu/hst/wfc3}:\\ two E2V thinned, backside illuminated and UV optimized 2k$\\times$4k CCDs contiguous on the long side of the chip, and covering a field of view (FoV) of $\\sim160$$\\times$$160$ arcsec$^2$.  The $\\omega$~Cen data set used here consists of 9$\\times$350~s exposures in each of the filters F225W, F275W, and F336W. The nine images follow a squared 3$\\times$3 dither-pattern with a step of about 40 arcsec (i.e.\\ $\\sim$1000 pixels), and were all collected on July 15, 2009.  We downloaded the standard pipe-line reduced {\\sf FLT} files from the archive. The {\\sf FLT} images are de-biased and flat-field corrected, but {\\it no} pixel-resampling is performed on them.  The {\\sf FLT} files are multi-extension fits (MEF) on which the first slot contains the image of what --hereafter-- we will call chip 1 (or simply [1]). The second chip, instead, is stored in the fourth slot of the MEF, and we will refer to it as chip 2 (or [2]).  [Note that others might choose a different notation].  Our GD corrections refer to the raw pixel coordinates of these images.  The fluxes and positions were obtained from a code mostly based on the software {\\sf img2xym\\_WFI} by Anderson et al.\\ (2006).  This is essentially a spatially variable PSF-fitting method.  We were pleased to see that for the WFC3/UVIS images of this data set the PSFs were only marginally undersampled.  Left panel of Figure~\\ref{fig:0A} shows a preliminary color-magnitude diagram in the three filters for the bright stars in the WFC3/UVIS data set.  In a future paper of this series we will discuss the PSF, its spatial variation and stability, as well as L-flats\\footnote{Residual low-frequency flat-field structure (L-flat) cannot be accurately determined from ground-based calibration data or internal lamp exposures.  L-flats need to be determined from on-orbit science data, for example from multiple observations of stellar fields with different pointings and roll angles (van der Marel 2003).}, pixel-area corrections, and recipes for deep photometry in stacked-images.  \\begin{figure*}[t!] \\centering \\includegraphics[width=14cm]{fig01r.ps} \\caption{ {\\it Left:} Preliminary color-magnitude diagram of the bright stars in the new WFC3/UVIS data set (fluxes are neither pixel-area- nor L-flat-corrected).  {\\it Right:} Color-magnitude diagram of the stars in our ACS/WFC master frame (from Villanova et al.\\ 2007).  Both plots are in instrumental magnitudes.  } \\label{fig:0A} \\end{figure*}  \\begin{figure*}[t!] \\centering \\includegraphics[width=14cm]{fig02.ps} \\caption{{\\it Left:} Color-magnitude diagrams of the high S/N stars in common with the ACS/WFC master frame (F435W), actually used to derive the geometric distortion correction, for each of the WFC3/UVIS filters (F225W, F275W, and F336W). } \\label{fig:0B} \\end{figure*}   ~\\\\ \\newpage ~\\\\ \\newpage ~\\\\ \\newpage ",
        "conclusions": "\\label{sec:3}  By using a limited (but best available) number of exposures with large dithers, and an existing ACS/WFC astrometric flat field, we have found a set of third-order correction coefficients to represent the geometric distortion of WFC3/UVIS in three broad-band ultraviolet filters. The solution was derived independently for each of its two CCDs.  The use of these corrections removes the distortion over the entire area of each chip to an average accuracy of $\\sim$0.025 pixel (i.e.\\ $\\sim$1 mas), the largest systematics being located in the $\\sim$200 pixels closest to the boundaries of the detectors (and never exceeding 0.06 pixels).  We advise the use of the inner parts of the detectors for high-precision astrometry.  The limitation that has prevented us from removing the distortion at an even higher level of accuracy is the lack of enough observations collected at different roll-angles and dithers which could enable us to perform an auto-calibration.  Nevertheless, the comparison of the mid-2002 ACS/WFC positions with the new WFC3 observations corrected with our astrometric solutions are good enough to clearly show the internal motions of $\\omega$~Centauri. These proved to be in perfect agreement with the most recent determinations.  We also derived the average absolute scale of the detector with an accuracy limited by the uncertainties in the plate-scale variations induced by the velocity aberration of the telescope motion in the Earth-Sun system.  For the future, more data with a longer time-baseline are needed to better characterize the GD stability of {\\it HST} WFC3/UVIS detectors in the medium and long term."
    },
    "0910/0910.1580_arXiv.txt": {
        "abstract": "We present studies of the collapse of neutron stars that undergo a hadron-quark phase transition. A spherical Lagrangian hydrodynamic code has been written. As initial condition we take different neutron star configurations taking into account its density, energy density and pressure distribution. The phase transition is imposed at different evolution times. We have found that a significant amount of matter on the surface can be ejected while the remaining star rings in the fundamental and first pressure modes. ",
        "introduction": "Black holes ($BHs$) and neutron stars ($NSs$) are certainly two major potential sources of gravitational waves ($GWs$). Unlike $BHs$, whose gravitational waveforms are specified essentially by their masses and angular momenta, the characteristics of the gravitational emission from $NSs$ depend on the properties of the nuclear matter.  When nuclear matter is compressed to a sufficiently high density, it turns into uniform three-flavor ($u$, $d$, and $s$) strange quark matter (SQM), since it is expected that SQM may be more stable than nuclear matter. The deconfined quark matter appears when the density is so high that the nucleons are ``touching'' each other. At this point, when the number density of nucleons $n\\sim O(1\\ {\\rm fm}^{-3})$, the quarks lose their correlation with individual nucleons and appear in a deconfined phase. Since the density required for this to happen is not much higher than nuclear matter density ($0.16\\ {\\rm fm}^{-3}$), the dense cores of neutron stars are the most likely places where the phase transition to quark matter may occur astrophysically. It should be noted that strange quarks (in a confined phase) could already exist in neutron stars with a hyperon core\\cite{marranghello}.  In principle, the existence of a thin crust of normal matter is possible at the surface of a strange star. On the other hand, if SQM is metastable at zero pressure, so that it is relatively more stable than nuclear matter only because of the high pressure in the cores of neutron stars, then the final products of the phase transition would be hybrid stars, which consist of quark matter cores surrounded by normal matter outer parts.  The above arguments were made in \\cite{lin} and those authors used polytropic equation of state to describe the complex nuclear matter structure. We proceed similar calculations with a realistic equation of state based on a field theoretical model. We also make use of a spherical Lagrangian hydrodynamical code instead of the Eulerian model used in \\cite{lin}. ",
        "conclusions": "\\indent We have obtained some important results which seems to be in complete agreement with those obtained by \\cite{lin}. The equation of state is shown in Fig.1 for three different parameter $\\alpha$ which leads to different hadron-quark phase transition pressure and density. Three stages of evolution are shown in Fig.2. The star radius shrinks and the central density is increasing during the collapse. In Fig.3 one can find the density evolution of each shell. The oscillation of the inner shells is clear in the final stage of evolution as it is the ejection of the external shells. The fourier transform is shown in Fig.4. This result shows a peak very close to the one shown by \\cite{lin}, at 3.1kHz, which shall correspond to the fundamental mode as well as other peaks, corresponding to the first pressure modes. It is also easy to find that most of the energy is released in the fundamental mode, as always claimed. In Fig.5 the radius and density evolution are plotted against time in this same figure.  Finally, after a future detection, one can try to fit the nuclear matter parameters to obtain the frequency of the detection, remembering that, if the frequency falls on the 3.0-3.4kHz bandwidth, the Brazilian Schenberg antenna will be able to identify the direction of the source, as shown in \\cite{mara}.  \\ack{We would like to thank CNPq and Fapesp for financial support.}"
    },
    "0910/0910.2126_arXiv.txt": {
        "abstract": "The cosmic evolution of the metal content of the intergalactic medium puts stringent  constraints on the properties of galactic outflows and on the nature of UV background. In this paper, we present a new measure of the redshift evolution of the mass density of \\CIV, $\\Omega_{\\rm CIV}$, in the interval $1.5 \\la z \\la 4$ based on a sample of  more than 1500 \\CIV\\ lines with column densities $10^{12} \\la N($\\CIV$) \\la 10^{15}$ cm$^{-2}$. This sample more than doubles the absorption redshift path covered in the range $z<2.5$ by previous samples. The result shows a significant increase of $\\Omega_{\\rm CIV}$ towards the lower redshifts at variance with the previously pictured constant behaviour. ",
        "introduction": "The cosmological mass density of \\CIV, $\\Omega_{\\rm CIV}$, observed as a function of redshift is a fundamental quantity closely related to the metal enrichment of the intergalactic medium (IGM). Its apparent lack of evolution  in the redshift interval $z\\simeq [1.5,5]$ \\citep{songaila01,pettini03,BSR03} is puzzling since both the physical conditions of the IGM and the properties of the ionizing background are thought to evolve between these epochs.  Remarkable efforts have been spent in recent years to extend the measure of  $\\Omega_{\\rm CIV}$ to redshift larger than 5 \\citep{ryanweber06,simcoe06} where a decrease of the star formation rate density is observed \\citep{bunker}. If $\\Omega_{\\rm CIV}$ is dominated by the metals produced in situ by the observed star forming galaxies, we would expect a decrease of its value at those redshifts. Vice versa, the value of $\\Omega_{\\rm CIV}$ could remain constant if it reflects the metallicity of a diffuse medium pre-enriched at very high redshift.  It should be noted, however, that this is a simplified scenario since, as redshift increases, the observed \\CIV\\ absorptions likely trace gas in structures of decreasing over-density and also the ionizing spectrum evolves in shape and intensity. As a consequence, the behaviour of $\\Omega_{\\rm CIV}$ could be different from that of $\\Omega_{\\rm C}$ and of the mean IGM metallicity \\citep[see e.g.][]{schaye03}.  The most recent measurements of \\CIV\\ absorptions in spectra of QSOs at $z\\sim6$ seem to indicate a downturn in the \\CIV\\ mass density at $z>5$ \\citep{becker09,ryanweber09}, though based only on 3 detected \\CIV\\ lines.  At redshift $z\\la 4.5$, a fundamental measurement of $\\Omega_{\\rm CIV}$  has been carried out by \\citet[][S01]{songaila01}. However, the redshift interval $1.5<z<2$ is poorly sampled by the considered QSO spectra. A more uniform redshift coverage is provided by the sample of \\citet{BSR03} although with fewer QSO spectra. Both analysis are consistent with a constant behaviour of $\\Omega_{\\rm CIV}$ in the range $[1.5,4.5]$. At $z < 1$, recent results based on HST UV data \\citep{cooksey} give $\\Omega_{\\rm CIV} = (6\\pm1) \\times 10^{-8}$ corresponding to a $2.8\\pm0.5$ increase over the $1.5 < z < 5$ values.  In this paper, we present a new measurement of  $\\Omega_{\\rm CIV}$ in the redshift range $[1.5,4]$ based on a sample of 25 high resolution, high signal-to-noise QSO spectra plus an additional sample of 8 QSO spectra from the literature.  The rest of the paper is organized as follows. The data are presented in \\S~2. In \\S~3 the analysis is carried out with the computation of $\\Omega_{\\rm CIV}$. The results are discussed in \\S~4. Throughout this paper, we assume $\\Omega_{\\rm m} = 0.26$, $\\Omega_{\\Lambda} = 0.74$ and $h \\equiv H_0/(100 {\\rm km\\ s}^{-1} {\\rm Mpc}^{-1}) =0.72$. ",
        "conclusions": "In this paper, we have presented a new determination of the cosmological mass density of \\CIV, $\\Omega_{\\rm CIV}$, at $z=[1.5,4]$ based on a large sample of high resolution, high signal-to-noise QSO spectra which more than doubles, with respect to previous measurements, the covered absorption path for $1.5 \\la z \\la 2.5$. The main result of our calculation is that  $\\Omega_{\\rm CIV}$ is no longer approximately constant in the considered redshift range, but shows a steady increase from $z\\sim3-5$ to $z\\sim1.5-2$. On the other hand, it appears that the \\CIV\\ mass density is not evolving significantly from these redshifts down to the present epoch \\citep{cooksey}.  The value of $\\Omega_{\\rm CIV}$ for the column density interval $12 \\leq \\log N($\\CIV$) \\leq 15$ can be converted into the IGM carbon content by mass, with the formula: \\begin{equation} Z_{\\rm C} \\simeq \\frac{\\Omega_{\\rm CIV}}{\\Omega_{\\rm b}}\\cdot \\frac{{\\rm C}}{{\\rm C\\,IV}} \\end{equation} where, $\\Omega_{\\rm b} = 0.0224/h^2$ \\citep{pettini08} is the contribution of baryons to the critical density and \\CIV/C is the fraction of C which is triply ionized, which depends on the assumed ionizing background. Assuming the maximum relative abundance of \\CIV, that is \\CIV/C~$\\la 0.5$, and the solar carbon over hydrogen abundance by mass $Z_{\\rm C,\\sun}= 0.0029 $ \\citep{asplund+05}, we obtain a lower limit at redshifts $[2.5,4.0]$ of $Z_{\\rm C} \\ga 1.4 \\times 10^{-6} = 4.9 \\times 10^{-4} Z_{\\rm C,\\sun}$ or [C/H]~$\\ga -3.3$. At redshift $[1.5,2.0]$: $Z_{\\rm C} \\ga 4.6 \\times 10^{-6} = 1.6 \\times 10^{-3} Z_{\\rm C,\\sun}$ or [C/H]~$\\ga -2.8$, corresponding to an increase of a factor of $\\sim 3$ with respect to the high redshift bin.  The metallicity measured at high redshift  is in agreement with the result obtained at $z=3$ by \\citet{schaye03} with the pixel optical depth method, in the case of over-densities above $\\sim 10$.  This is consistent with the fact that $\\Omega_{\\rm CIV}$ traces mainly the metallicity of the over-dense regions in the proximity of galaxies and its redshift evolution is linked with that of the strong systems. \\citet{simcoe04} found a median metallicity [C,O/H]~$\\simeq -2.8$ for the intergalactic gas at $z\\sim 2.5$ with a method based on the direct detections of the absorption lines which takes into account also the upper limits. Re-normalizing our result to the solar abundances and the fraction of C triply ionized adopted by those authors  (\\CIV/C~$\\la 0.25$), we obtain  [C/H]~$\\ga -3.0$ consistent with their result.   The physical explanation for the redshift evolution of $\\Omega_{\\rm CIV}$ can be explored with cosmological simulations. We will devote a subsequent paper to the comparison between observations and predictions of hydro-simulations with a detailed treatment of metal enrichment (Tescari et al., in preparation). In the following, we will briefly discuss the few predictions of $\\Omega_{\\rm CIV}$ present in the literature and compare them with our results.  \\citet[][OD06]{opp_dave06} have run cosmological hydro-simulations of galaxy formation meant to reproduce the metallicity of the IGM, where metals are expelled from galaxies by momentum-driven winds (whose velocity is proportional to the dispersion velocity of the galaxies). Those winds are very effective in transporting the gas without heating it too much. Their fiducial simulation reproduces consistently the star formation rate of the Universe, the volume averaged metallicity of the IGM and the lack of evolution of   $\\Omega_{\\rm CIV}$ from $z\\approx 5 \\rightarrow 1.5$, as observed in the data available at the time. On the other hand, the predicted mass density of C increases by nearly an order of magnitude toward lower redshifts in the same redshift interval, due to the increase in the total metallicity of the gas. The constant behaviour of $\\Omega_{\\rm CIV}$ is obtained by balancing the increase in C abundance with  galactic winds. This form of feedback heats the IGM causing a decrease of the \\CIV\\ ionization fraction. In particular, the peak of the  \\CIV\\ ionization fraction distribution as a function of over-density shifts toward larger over-densities as redshift decreases. This is consistent with the increase in the ratio between strong and weak \\CIV\\ systems at lower redshifts shown in Fig.~\\ref{ndens} and the identification of strong systems with denser environment as deduced from the observed 2-point correlation function (see Section~3.3).  To conclude, none of the models proposed by OD06 predicts an increase of $\\Omega_{\\rm CIV}$ toward lower redshifts. An improvement in the physics of the wind, plus the introduction of the AGB feedback, results in a slight increase of $\\Omega_{\\rm CIV}$ in the ranges $z < 0.5$ and $z > 5$ leaving substantially unchanged the other values \\citep[][OD08]{opp_dave08}.    An effect which has not been taken into account up to now is the evolution in the shape of the UV background due to the \\HeII\\ re-ionization process expected at redshift $\\sim 3$. Naively, due to the increase in the number of free hard photons (ionizing \\CIV\\ into \\CV) at the end of the re-ionization epoch, we would expect a trend opposite to what is observed: a decrease of the amount of \\CIV\\ after $z\\sim3$.  However, the details of this process are still under study \\citep[e.g.][]{bolton09,hm09} and other factors should be accounted for properly: the decrease in the number density of QSOs going towards lower redshifts and the fact that at low $z$ the regions traced by the strong \\CIV\\ absorbers (driving the evolution of  $\\Omega_{\\rm CIV}$) could be only mildly affected by the cosmic UV background.  The observed raise of the cosmic \\CIV\\ mass density in the redshift range $1.5-2.5$ puts a strong constraints on the models describing the interplay between galaxies and their surrounding medium, suggesting that something is still missing in the physical implementation of galactic feedback."
    },
    "0910/0910.5706_arXiv.txt": {
        "abstract": "We discuss current cosmological constraints on axions, as well as future sensitivities. Bounds on axion hot dark matter are discussed first, and subsequently we discuss both current and future sensitivity to models in which axions play the role as cold dark matter, but where the Peccei-Quinn symmetry is not restored during reheating. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.2310_arXiv.txt": {
        "abstract": "We present an abundance analysis of the star Cernis~52 in whose spectrum we recently reported the napthalene cation in absorption at 6707.4 {\\AA}. This star is on a line of sight to the Perseus molecular complex. The analysis of high-resolution spectra using a $\\chi2$-minimization procedure and a grid of synthetic spectra provides the stellar parameters and the abundances of O, Mg, Si, S, Ca, and Fe. The stellar parameters of this star are found to be $T_{\\mathrm{eff}} = 8350 \\pm 200$ K, \\loggl $= 4.2 \\pm 0.4$ dex. We derived a metallicity of $\\mathrm{[Fe/H]} = -0.01 \\pm 0.15$. These stellar parameters are consistent with a star of $\\sim 2$~\\Msun in a pre-main-sequence evolutionary stage. The stellar spectrum is significantly veiled in the spectral range $\\lambda\\lambda5150-6730$~{\\AA} up to almost 55 per cent of the total flux at 5150~{\\AA} and decreasing towards longer wavelengths. Using Johnson-Cousins and 2MASS photometric data, we determine a distance to Cernis~52 of 231$^{+135}_{-85}$~pc considering the error bars of the stellar parameters. This determination places the star at a similar distance to the young cluster IC~348. This together with its radial velocity, $v_r=13.7\\pm1$ \\kmso, its proper motion and probable young age support Cernis~52 as a likely member of IC~348. We determine a rotational velocity of $v\\sin i=65 \\pm 5$~\\kms for this star. We confirm that the stellar resonance line of \\ion{Li}{1} at 6707.8~{\\AA} is unable to fit the broad feature at 6707.4~{\\AA}. This feature should have a interstellar origin and could possibly form in the dark cloud L1470 surrounding all the cluster IC~348 at about the same distance. ",
        "introduction": "In a recent study of the diffuse interstellar bands (DIBs) toward a region of anomalous microwave emission, Iglesias-Groth et al. (2008) have shown evidence for the presence of the naphthalene cation, ${\\rm C}_{10}{\\rm H}_8^+$, the simplest Policyclic Aromatic Hydrocarbon (PAH), in the line of sight towards the moderately reddened star Cernis~52 (BD+31$^o$ 640, $V=11.4$, $E(B-V)=0.9$, Cernis 1993). The PAHs, including the naphthalene cation, have been also detected in a cometary dust sample returned to Earth by Stardust and in interplanetary dust particles, possibly of cometary origin (see the review by Li 2009).  Cernis~52 is an early type star, classified as A3V by Cernis (1993), has coordinates $\\alpha = 03^{\\rm h}43^{\\rm m}00.3^{\\rm s}$ and $\\delta = +31^\\circ58'26''$ (J2000.0), and is located at an angular separation of less than one degree from the very young stellar cluster IC~348 in the Perseus OB2 molecular complex. The photometric distance to star Cernis~52 is 240 pc (Cernis 1993), consistent with a location in the molecular complex OB2 where the microwave emitting cloud is also most likely located (see Watson et al. 2005). The uncertainties prevent to conclude whether the star is embedded in a cloud of this young star forming region or lay behind.  Iglesias-Groth et al. (2008) claimed that the relative strength of several diffuse interstellar bands detected in the spectrum of Cernis~52 was consistent with the presence of a molecular cloud in the line of sight. The anomalous microwave emission detected in the direction to Perseus (Watson et al. 2005, 2006) could be associated to electric dipole radiation of fast spinning hydrogenated carbon-based molecules in a molecular cloud, a mechanism originally proposed by Draine and Lazarian (1998). It is important to establish whether the cloud responsible for the excess color of Cernis~52 hosts the carriers of both, the diffuse interstellar bands observed in the spectrum of this star and the anomalous microwave emission detected in its line of sight.  Iglesias-Groth et al. (2008) also noted the near-coincidence in wavelength between the interstellar cation's feature and the stellar Li I resonance line. To gain insight into the contribution of the stellar line to the observed broad feature, we undertook a thorough analysis by spectrum synthesis of Cernis 52's spectrum with two principal goals - a general abundance analysis for the star and spectrum synthesis fits to the spectrum around 6707~{\\AA} in order assess the Li I line's contribution.  Here we present a precise determination of the stellar parameters and a detailed chemical composition study of the star Cernis 52. This abundance analysis should provide indications of the likely Li abundance for the star, i.e., if the star is not chemically peculiar should have a lithium abundance no greater than the cosmic Li abundance of A(Li)\\footnote{A(Li)$=\\log [N({\\rm Li})/N({\\rm H})]+12$}\\,$=3.3.$ In this paper we will discuss with special attention the naphthalene cation's feature at 6707.4~{\\AA} in the spectrum of Cernis~52 and investigate the relationship of this star with the Perseus star forming region. ",
        "conclusions": "We have performed a detailed chemical analysis of the star Cernis~52. We apply a technique that provides a determination of the stellar parameters, taking into account any possible source of veiling. We find $T_{\\mathrm{eff}} = 8350 \\pm 200$ K, $\\log (g/{\\rm cm~s}^2) = 4.2 \\pm 0.4$, $\\mathrm{[Fe/H]} = -0.01 \\pm 0.15$, and a veiling (defined as $F_{\\rm  veil}/F_{\\rm total}$) of less than 55\\% at 5000 {\\AA} and decreasing toward longer wavelengths.  The spectrum of Cernis~52 shows many features very likely related with the interstellar medium. In addition, we discover in photometric images a companion 1.7~mag fainter star at a distance of $0.818\\arcsec\\pm0.007\\arcsec$, but this star only contributes with 10\\% of the stellar flux in our spectra and hardly affect the stellar features of Cernis~52. The derived chemical abundances are roughly solar within their error bars. This prevent the star from being a chemically peculiar star.  We have determined the radial velocity of Cernis~52 at $v_r=+13.7 \\pm 1$~\\kmso, being almost equal to the mean radial velocity of the young cluster IC~348.  We have compared the stellar parameters with pre-main-sequence evolutionary tracks of solar metallicity and see that the star is consistent with being a pre-main-sequence A-type star with an age of $3-20$~Myr.  We have estimated the distance to Cernis~52 using the available Johnson-Cousins and 2MASS photometric data, at 231$^{+135}_{-85}$~pc according to the stellar parameters and its error bars. This value also agrees with the distance to the cluster IC~348.  The proper motion of Cernis~52, $(\\mu_\\alpha \\cos \\delta$, $\\mu_\\delta=(+7.22$,$-8.62)\\pm(1.5$,$1.6)$ mas~yr$^{-1}$, is consistent with the proper motion of IC~348.  All these measurements make it likely that the star Cernis~52 (\\mbox{BD+31$^o$ 640}) belongs to the young cluster IC~348.  We confirm that the feature at 6707.4~{\\AA} is not related with a stellar lithium line because the rotational velocity of the star, $v\\sin i=65 \\pm 5$~\\kmso, is too low to explain the broad feature associated with the naphthalene cation. Furthermore, the presence of a companion star cannot either explain this feature even for an abnormally high Li abundance, $\\log [N({\\rm Li})/N({\\rm H})]+12 > 5$ dex.  As already stated in Cernis (1993), the interstellar features, in particular, the naphthalene cation, that appear in the spectrum of Cernis~52 may form in the dark cloud L1470 which covers all the cluster IC~348 and is at about the same distance."
    },
    "0910/0910.5476_arXiv.txt": {
        "abstract": "We present {\\it Spitzer} observations of a sample of 12 starless cores selected to have prominent 24~\\micron\\ shadows.  The {\\it Spitzer} images show 8 and 24~\\micron\\ shadows and in some cases 70~\\micron\\ shadows; these spatially resolved absorption features trace the densest regions of the cores.  We have carried out a $^{12}$CO (2-1) and $^{13}$CO (2-1) mapping survey of these cores with the Heinrich Hertz Telescope (HHT).  We use the shadow features to derive optical depth maps.  We derive molecular masses for the cores and the surrounding environment; we find that the 24~\\micron\\ shadow masses are always greater than or equal to the molecular masses derived in the same region, a discrepancy likely caused by CO freeze--out onto dust grains.  We combine this sample with two additional cores that we studied previously to bring the total sample to 14 cores.  Using a simple Jeans mass criterion we find that $\\sim\\!2/3$ of the cores selected to have prominent 24~\\micron\\ shadows are collapsing or near collapse, a result that is supported by millimeter line observations.  Of this subset at least half have indications of 70~\\micron\\ shadows.  All cores observed to produce absorption features at 70~\\micron\\ are close to collapse.  We conclude that 24~\\micron\\ shadows, and even more so the 70~\\micron\\ ones, are useful markers of cloud cores that are approaching collapse. ",
        "introduction": "Stars are born in cold cloud cores \\citep[e.g.,][]{bergin07}, where gas and dust are compressed to high enough densities to cause the core to start collapsing. One of the most critical questions in astronomy is what the physical conditions for core formation and collapse are. Therefore, observations of cores in the rare stage during, or just prior to, collapse are critical to constrain theories of the very earliest stages of star-formation \\citep[e.g.,][]{shu87,ballesteros03,myers05}.  The initial process of collapse, or the transition from a stable core to a core with an embedded protostar, happens rapidly \\citep{hayashi66}; additionally, the cores are necessarily dense and cold. These conditions pose significant observational challenges to identify cores close to collapse. Because of the temperature ranges involved, $\\sim\\!10$~K \\citep[e.g.,][]{lemme96,caselli99,hotzel02}, and high densities, $\\sim$10$^5$ cm$^{-3}$ \\citep[e.g.,][]{bacmann00}, cold cloud cores can best be observed at far-infrared, submillimeter, and millimeter wavelengths \\citep[e.g.,][]{stutz09}.  Such observations of dense cores show that most of them are close to equilibrium and not collapsing \\citep[e.g.,][]{lada08}. Furthermore, if the cores are not supported by thermal and turbulent pressure alone, then a modest magnetic field can halt collapse \\citep[e.g.,][]{kandori05,stutz07,alves08}.  Millimeter line observations are the traditional way to search for collapsing cores \\citep[e.g.,][]{walker86}; however, there are ambiguities in the interpretation of such measurements \\citep[e.g.,][]{walker88,menten87,mundy90,narayanan98}. Here we use mid-infrared shadows as an alternative observational approach to the study of starless cores, one that does not depend on the interpretation of millimeter line profiles and how these profiles relate to the underlying velocity field of the core material. Additionally, our method is sensitive to higher column densities than traditional studies of starless cores by background star extinction, usually limited to A$_V$ $\\le$ 30 magnitudes \\citep[e.g.,][]{alves01,kandori05}.  Cores that are isolated from regions of massive star formation are very useful to understand the initial stages of collapse into stars. These cores are free from complicating factors, e.g., massive core fragmentation, feed-back from high-mass stars, and protostellar outflows \\citep[e.g.,][]{lada03,dewit05,fallscheer09}, to name a few, that can have large effects on a study aimed at understanding how individual low mass stars form. We present a sample of 12 relatively isolated cores that were selected from {\\it Spitzer} MIPS observations of Bok globules and other star-forming regions to have prominent 24 $\\mu$m shadows, spatially resolved absorption features caused by the dense core material viewed in absorption against the interstellar radiation field. We also include 2 previously studied cores, CB190 \\citep{stutz07} and L429 \\citep{stutz09}, for a final sample of 14 cores.  Such objects with 24 $\\mu$m shadows always show a counterpart 8 $\\mu$m shadow, analogous to the IR absorption features produced by more distant and massive structures termed infrared dark clouds \\citep[IRDCs; e.g.,][]{perault96,butler09,peretto09,lee09} and the ISO 7 $\\mu$m absorption features studied by \\citet{bacmann00}. We also present Heinrich Hertz Telescope (HHT) $^{12}$CO (2-1) and $^{13}$CO (2-1) on-the-fly (OTF) maps of regions $\\sim$ 10$'$ on a side surrounding each core.  Our search method is efficient at identifying cores that are near collapse.  In \\S~2 we describe the observations and data processing; in \\S~3 we derive optical depth maps from the 8 and 24 $\\mu$m images; in \\S~4 we present our mass measurements and we describe the stability analysis that we apply to the sample of cores; and finally, in \\S 5 we summarize our main conclusions. ",
        "conclusions": "We study the 8~\\micron, 24~\\micron, and 70~\\micron\\ shadows cast by a sample of 14 starless cores.  We derive 24~\\micron\\ core masses and sizes; we apply a Jeans mass criterion, and attempt to account for turbulent and magnetic support in the cores, in order to assess the collapse state of each core.  We caution that distance uncertainties can have a large effect on our Jeans mass analysis.  In addition, we have obtained $^{12}$CO (2-1) and $^{13}$CO (2-1) OTF maps of the cores; the molecular core masses we derive are always less than the 24~\\micron\\ masses.  This discrepancy is likely caused by freezeout onto dust grains.  From this work we conclude that:  \\noindent 1. 70\\% (10/14) of cores selected to have prominent 24~\\micron\\ shadows seem to be approaching collapse, indicating that this criterion selects dense, evolved cores.  \\noindent 2. 50\\% (5/10) of cores that are classified as approaching collapse have indications of 70~\\micron\\ shadows; the 70~\\micron\\ data quality does not allow for a rigorous analysis of these shadows.  \\noindent 3. All cores with indications of 70~\\micron\\ shadows have millimeter line profiles showing blue asymmetries, indicating that these long--wavelength shadow features are produced by very evolved cores.  These cores are all likely to be close to collapse.  \\noindent 4. Shadows at 24~\\micron\\ and especially at 70~\\micron\\ appear to be effective indicators of cores that are approaching collapse."
    },
    "0910/0910.0854_arXiv.txt": {
        "abstract": "The Minimal Supersymmetric Standard Model has several flat directions, which can naturally be excited during inflation. If they have a slow (perturbative) decay, they may affect the thermalization of the inflaton decay products. In the present paper, we consider the system of udd and QLd flat directions, which breaks the $U(1)\\times SU(2)\\times SU(3)$  symmetry completely. In the unitary gauge and assuming a general soft breaking mass configuration, we show that for a range of parameters, the background condensate of flat directions can undergo a fast non-perturbative decay, due to non-adiabatic evolution of the eigenstates. We find that both the background evolution and part of the decay can be described accurately by previously studied gauged toy models of flat direction decay. ",
        "introduction": "\\label{sec:intro} Flat directions are generic features of supersymmetric theories. They are directions in field space along which the renormalizable part of the scalar potential vanishes. The Minimal Supersymmetric Standard Model (MSSM) and its extensions have a plethora of D and F-flat directions \\cite{Gherghetta:1995dv},  which are lifted due to supersymmetry breaking. During inflation, if their effective mass is small compared to the Hubble rate, the fields can develop large vacuum expectation values (VEV) along the flat directions of the potential \\cite{developvev,Ellis:1987rw}. This growth is bounded above by the non-renormalizable term with lowest dimension, which has the form $\\phi^d/M^{d-3}$ with $d \\geq 4$. If inflation is long enough, the growth will proceed up to $\\langle \\phi \\rangle \\sim (m_\\phi M^{d-3})^{1/(d-2)}$ \\cite{Dine:1995kz}. If all non-renormalizable terms allowed by gauge invariance are present, each class of flat direction will be lifted by the term with the smallest $d$ allowed by MSSM symmetries \\cite{Gherghetta:1995dv}. On the other hand, discrete symmetries may forbid some of such terms and a non-renormalizable term with a higher $d$ may determine the VEV of the flat direction. All possible flat directions in MSSM are excited by terms with $d \\le 9$ \\footnote{The term at which all flat directions are lifted may be different for the extensions of MSSM. For example, in $\\nu$MSSM, no flat direction survives beyond $d=6$ \\cite{Basboll:2009tz}.}.  The formed condensate can have several cosmological implications: In the presence of phase dependent potential terms, the flat directions may source a finite baryon number density through the Affleck-Dine mechanism \\cite{Affleck:1984fy, Linde:1985gh, Allahverdi:2000zd}. It has also been suggested that they may be responsible for inflation \\cite{Allahverdi:2006iq}. Additionally, due to the large VEV of the flat directions, all the fields coupled to them acquire a large effective mass, slowing down the decays they mediate and resulting in a small perturbative decay rate.  Typically, the perturbative decay of flat direction concludes after $\\sim 10^{11}$ rotations \\cite{Olive:2006uw}. These long lived flat directions also keep the gauge fields of broken symmetries (assumed to be all the Standard Model ones) heavy, suppressing the scatterings among the inflaton decay products, thus delaying their thermalization \\cite{Allahverdi:2005mz}. In addition, the energy density of the (relativistic) inflaton decay products may become sub-dominant over that of (massive) flat directions. The subsequent radiation stage will then be dominated by the thermal distribution of flat direction decay products, rather than those of the inflaton. These effects on thermalization require sufficiently large initial flat direction VEVs, which can be acquired only if non-renormalizable superpotential terms up to $d=11$ are absent \\cite{Olive:2006uw}.  However, if the decay of the flat directions is controlled by non-perturbative effects, the effect on thermalization will be very different than the above picture. This possibility was first discussed in \\cite{Allahverdi:1999je}, in the framework of a toy model based on F-term type interactions. For this model, the frequencies of the particles coupled to the flat directions evolve adiabatically, not allowing a resonant decay. On the other hand, it was shown in \\cite{Olive:2006uw}, that the D-term potential provides non-trivial interactions among the perturbations through a non-diagonal and time dependent mass matrix. Even if the eigenvalues of this matrix evolve adiabatically, the diagonalization procedure itself may be non-adiabatic, due to a fast rotation of the eigenvectors. The resulting exponential decay of the condensate has a much higher rate than the perturbative one, giving a decay after $\\mathcal{O}(10)$ rotations of the flat directions. In \\cite{Olive:2006uw}, it was also argued that at least two or more flat directions need to be excited for this effect to be realized. The argument is as follows: Since the resonant effect occurs in the D-terms, only the perturbations coupled to the VEVs through the symmetry generators are counted. Out of these degrees of freedom, two per broken symmetry will correspond to a Higgs and a Goldstone. Furthermore, two more (light) degrees of freedom will decouple, corresponding to the real and imaginary parts of fluctuations along each flat direction. In order to have a non-adiabatic mixing, one needs additional light degrees of freedom that the condensate can decay into. To formulate, the number of remaining degrees of freedom present in the system will be \\begin{equation} \\left( \\begin{array}{c} {\\rm remaining}\\\\ {\\rm degrees} \\end{array} \\right) = \\left( \\begin{array}{c} {\\rm d.o.f.} \\\\ {\\rm in~D~terms} \\end{array} \\right) - 2\\,\\times \\left( \\begin{array}{c} {\\rm broken} \\\\ {\\rm symmetries} \\end{array} \\right)-2\\,\\times \\left( \\begin{array}{c} {\\rm flat} \\\\ {\\rm directions} \\end{array} \\right)\\,. \\label{counting} \\end{equation} As long as this number is zero, there will not be any room for non-perturbative decay. For instance, for the typical cases of single flat directions, no residual degree of freedom is present \\cite{Olive:2006uw}. On the other hand, there exist flat directions that are non-exclusive, i.e. they do not give a large mass to each other due to their VEVs. If the conditions to excite a single flat direction are present, one can expect that the whole set of flat directions non-mutually exclusive with that one is excited. If realized, such a case would provide the extra degrees of freedom into which the condensate may decay non-perturbatively.  The longevity of single flat directions was later reiterated by the authors of \\cite{Allahverdi:2006xh}, where it was also argued that for the non-perturbative decay of multiple flat directions, one needs some degree of tuning of the initial VEVs: Since different flat directions may be lifted by different non-renormalizable terms in the superpotential \\cite{Gherghetta:1995dv}, one may in general expect hierarchical VEVs. Such a case reduces to a single flat direction, which decays only perturbatively. The maximum amount of hierarchy that can provide a non-perturbative decay depends on the ellipticity of the orbits of the VEVs in their complex plane. In later works, gauged toy models with two flat directions \\cite{Basboll:2007vt,Gumrukcuoglu:2008fk} and examples from MSSM \\cite{Basboll:2008gc} were studied, each verifying that multiple flat directions may decay non-perturbatively. Additionally, in \\cite{Gumrukcuoglu:2008fk}, the fast decay was shown to be realized for a range of VEV ratios of three orders of magnitude. This range was found to be a consequence of the phase dependent terms introduced in the fashion of \\cite{Affleck:1984fy}.  However, the question of whether the toy models provide a good description of MSSM flat directions needs to be answered. For example, in \\cite{Gumrukcuoglu:2008fk}, the gauged toy models of four fields with only $U(1)$ or $SU(2)$ charges have been studied, yet in MSSM, no such flat direction configuration is possible and  generically, for multiple flat directions, the field content has charges of all symmetries. Furthermore, the production of the remaining degrees of freedom (\\ref{counting}) may be suppressed if they acquire large masses through the F-terms. Therefore, the main goal of the present work is to find a concrete example from MSSM for which, the decay of the flat directions proceed analogously to the gauged toy model case. The flat directions in the models of \\cite{Gumrukcuoglu:2008fk} are decoupled at the background level, and each of the two VEVs evolve independently like two single flat directions. However, having independently evolving VEVs is not a requirement for non-perturbative decay. For instance, flat directions with coupled VEVs also have the necessary ingredients for decay \\cite{Basboll:2007vt, Basboll:2008gc}. The latter systems are much more complicated than the former ones and a precise answer requires extended numerical calculation. Our primary focus will be on the system of $u^cu^cd^c$ and $QLd^c$ flat directions and we will show that their decay can be described by the four field toy model of \\cite{Gumrukcuoglu:2008fk}.  Additionally, we will address some issues arising from the assumption that the fluctuations along the flat directions are decoupled from the other modes. For instance, for the $QLd^c+LLe^c$ system, Ref. \\cite{Basboll:2008gc} claimed that the Higgses and the flat direction perturbations have non-adiabatic mixings. On the other hand, for the toy models of \\cite{Gumrukcuoglu:2008fk}, it was shown that the these degrees of freedom indeed decouple from the rest of the action. However, the latter result is a consequence of the assumption that the fields in a given flat direction have equal masses, an assumption not generically applicable to MSSM fields.  If the fields have distinct masses, the flat direction perturbations are no longer decoupled from the Higgses. If these mixings are non-adiabatic, they may result in a non-perturbative decay, even if the counting (\\ref{counting}) leaves no extra degrees of freedom. For the models we consider, we show that these mixings have negligible effect and the flat direction perturbations decouple as described in \\cite{Olive:2006uw, Gumrukcuoglu:2008fk}. We will first generalize the single flat direction toy model of \\cite{Gumrukcuoglu:2008fk} to have arbitrary masses and verify that the flat direction does not decay non-perturbatively. The approximations and methods we adopt in this simple example will provide us the necessary tools for the background evolution of $u^c u^c d^c$ and $Q L d^c$ flat directions, which will also be studied with generic mass terms.  The paper is organized as follows. In Section \\ref{sec:mssm}, we discuss different classes of multiple flat directions in MSSM over some examples which allow for non-perturbative decay, and determine which example is most likely to be described by the gauged $4$-field toy model. In Section \\ref{sec:formalism}, we review the formalism for the calculations of this decay. In Section \\ref{sec:1fd}, we generalize the single flat direction toy model with two complex scalar fields and $U(1)$ gauge field, to include arbitrary soft masses. This provides a basis for non-degenerate mass calculation of a more complicated model, carried out in the next section. In Section \\ref{sec:2fd} we present a complete study of $u^cd^cd^c+Q L d^c$ flat directions in MSSM, with arbitrary soft masses, where we compare the final results to the ones for the $4$-field toy model in \\cite{Gumrukcuoglu:2008fk}. The results are summarized and discussed in Section \\ref{sec:disc}. Finally, we include the technical steps of the calculation in the appendices at the end. ",
        "conclusions": "\\label{sec:disc}  The primary aim of this paper was to see if the extensive numerical study carried out in the framework of $4$-field gauged toy model \\cite{Gumrukcuoglu:2008fk} can accurately describe a realistic case. We have shown that this toy model, which has only $U(1)$ symmetry, two mass parameters and no F-terms, provides a very good description of an example from MSSM, where both $u^cd^cd^c$ and $Q L d^c$ flat directions are present simultaneously. Specifically, both cases contain ``independent flat directions'', that is, the fields in one flat direction are decoupled from the others at the background level. As a result, the background equations of motion for the two models have the same form. Furthermore, we found that out of the $64$ real degrees of freedom in the problem, only the two decoupled actions $S_{s_{\\bar{1}},s_2, d_{\\bar{2}}}$ and $S_{b_{\\bar{1}},s_3, d_{\\bar{3}}}$, each consisting of 8 degrees, exhibit the non-adiabatic evolution of the eigenstates. At the limit of degenerate masses, that is, when the masses of the fields in a flat direction are identical, we found that $S_{b_{\\bar{1}},s_3, d_{\\bar{3}}}$ can be decomposed into two copies of the coupled system in the $4$ field gauged toy model. Therefore, it is safe to state that, in the degenerate mass case, the numerical results of the toy model provides an exact description of this part of the action. If the action $S_{s_{\\bar{1}},s_2, d_{\\bar{2}}}$ contributes to the non-perturbative decay as we expect, the resulting quanta will be different ones and their production will not effect the occupation numbers of generated $b_{\\bar{1}}$, $s_3$ and $d_{\\bar{3}}$ perturbations, as long as the linearized approximation holds.  Another focus of the present paper was to understand the effect of different soft supersymmetry breaking masses to the non-perturbative production. For an exactly D-flat direction, the fields in that direction have the same soft mass. Conversely, if all the fields have different masses, the D-term will be non-zero and proportional to the fourth power of mass differences. As a reference calculation, we showed that for a single flat direction, the flat direction perturbations still decouple at the leading order, thus verifying that there is no non-perturbative decay in this case. The effect of different masses in the $u^cd^cd^c + Q L d^c$ example had similar consequences for the flat direction perturbations. Furthermore, the four field toy model with degenerate masses still describes the subsystem $S_{b_{\\bar{1}},s_3, d_{\\bar{3}}}$ exactly when the condition (\\ref{deltamcom}) holds. Therefore, based on the results of \\cite{Gumrukcuoglu:2008fk}, we conclude that the flat directions $u^cd^cd^c + Q L d^c$ decay in $\\mathcal{O}(10)$ rotations, also for this case.  One issue about the example considered here is the hierarchy between the VEVs corresponding to each flat direction. The first non-renormalizable operator present for $u^cd^cd^c$ flat direction has $d=6$, whereas for $Q L d^c$, it has $d=4$ \\cite{Gherghetta:1995dv}, resulting in an initial VEV ratio of $\\sim10^3$. This ratio, although large, may still give rise to a production if the orbits have enough ellipticity and the mass ratio is large enough. For instance, in the numerical results of \\cite{Gumrukcuoglu:2008fk}, it was shown that for two flat directions with mass ratio $\\tilde{m}/m =7.63$, they decay within $20$ rotations for initial VEV ratio of $10^3$. If the orbits are closer to the radial one, one might still have a non-perturbative decay with a smaller mass ratio for the same VEV hierarchy. On the other hand, the non-renormalizable terms are model dependent, and they may be forbidden by some discrete symmetries. In fact, to recover the conditions of delayed thermalization \\cite{Allahverdi:2005mz}, one needs $\\phi_{\\rm in} \\gtrsim 10^{-2} M_p$, which requires all terms with $d < 11$ to vanish \\cite{Olive:2006uw}.  Although the numerical results of \\cite{Gumrukcuoglu:2008fk} shows that the flat directions may decay non-perturbatively, it is still not clear how this effect changes the picture of thermalization. Specifically, the new quanta are produced in a resonant band with momenta $k \\lesssim m$ which is still non-relativistic. Since the variances are large, the gauge fields will still have large enough mass contributions to suppress the scatterings between inflaton decay products. On the other hand, tracking the evolution of the produced particle distribution is beyond the reach of the linearized calculation. One needs to control the back reaction effects to determine how fast the variances decrease. In such a computation, the distribution is likely to thermalize, possibly not in $\\mathcal{O}(10)$ rotations, but we expect it to be much earlier than $10^{11}$ rotations that is required by the perturbative decay \\footnote{Indeed, our preliminary results from lattice simulations of non-gauged toy model \\cite{Olive:2006uw} up to $\\mathcal{O}(100)$ rotations show the beginning of thermalization of the produced particles, i.e. the initial distribution with momenta $k \\lesssim m$ starts to extend toward high momentum region.}.  It is clear that the linear study done in the present paper, along with \\cite{Olive:2006uw, Basboll:2007vt, Basboll:2008gc, Gumrukcuoglu:2008fk} are limited to the stages of the evolution until the production is significant, and they only provide a glimpse to the beginning of the non-perturbative decay. For instance, we found that the $u^cd^cd^c + Q L d^c$ example should result in a decay within $\\mathcal{O}(10)$ rotations for the range of parameters in \\cite{Gumrukcuoglu:2008fk}. However, once non-linear effects become important, the particles produced through different decoupled actions may interact and change the outcome of this study. Therefore, one should be cautious to extrapolate a non-linear study based on the toy model to a realistic one.  The logical step to be taken next is to include higher order terms and study the effect on the decay time and thermalization of the produced quanta. In a recent study \\cite{Dufaux:2009wn}, the non-gauged toy model of \\cite{Olive:2006uw} was evolved on a lattice using the ClusterEasy code \\cite{Felder:2000hq}, verifying that the non-perturbative decay is still realized within the first few rotations. Another interesting result of \\cite{Dufaux:2009wn} was the calculation of the gravity waves, sourced by the quick decay of the flat directions. Their spectrum was found to fall naturally into Hz-kHz range and depending on the initial VEV of the flat directions, may potentially be within the reach of upcoming experiments, such as Advanced LIGO. For more realistic models with gauge fields, the resulting spectrum may be different, but since the mass scale of the flat directions is of order TeV, it will have a frequency range similar to that in the toy model. In a future work, we will address the effect of back-reaction and gravity wave production in the framework of gauged models.  However, there are still some problems left that can be dealt with analytical tools. In \\cite{CyrRacine:2009qk}, it was shown that the non-perturbative decay of flat directions may have an observable effect through the amplification of curvature perturbations. In the context of natural supergravity inflation \\cite{Kawasaki:2000ws}, Ref. \\cite{Kaminska:2009wh} showed that the non-perturbative decay of flat directions allows the inflaton preheating to be realized in these models. On the other hand, even at the linearized level, we do not have a complete numerical study of the decay for ``overlapping flat directions'', although we have approximate calculations showing the non-adiabatic rotation occurs \\cite{Basboll:2007vt, Basboll:2008gc}. Specifically, the toy model of \\cite{Basboll:2007vt} is qualitatively different than the ones considered in the present work, as well as \\cite{Gumrukcuoglu:2008fk}, in the sense that the flat directions are coupled at the background level. A numerical analysis in the fashion of \\cite{Gumrukcuoglu:2008fk} would be very useful in understanding the time scale of the decay and the range of hierarchy that allows production. There are many examples with ``overlapping flat directions'' in MSSM which provide the necessary ingredients for non-adiabatic evolution (e.g. \\cite{Basboll:2008gc}), and it is an interesting challenge to find simple toy models that correctly describe at least parts of these examples."
    },
    "0910/0910.2851_arXiv.txt": {
        "abstract": "name{Samenvatting}% \\def\\bibname{Bibliografie}\\def\\chaptername{Hoofdstuk}% \\def\\appendixname{Bijlage}\\def\\contentsname{Inhoudsopgave}% \\def\\listfigurename{Lijst van figuren}\\def\\listtablename{Lijst van tabellen}% \\def\\indexname{Index}\\def\\figurename{Figuur}\\def\\tablename{Tabel}% \\def\\partname{Deel}\\def\\enclname{Bijlage(n)}\\def\\ccname{Ter attentie van}% \\def\\headtoname{Aan}\\def\\headpagename{Pagina}% \\def\\today{\\number\\day\\space\\ifcase\\month\\or januari\\or februari\\or maart\\or% april\\or mei\\or juni\\or juli\\or augustus\\or september\\or oktober\\or% november\\or december\\fi \\space\\number\\year}% \\typeout{ >>>>> use hlatex209 for Dutch hyphenation <<<<< }} \\hyphenation{Schrij-ver Krij-ger Kuij-pers Bal-le-gooij-en}   \\def\\warningoverprint #1{\\special{!userdict begin /bop-hook{gsave 100 600 translate -45 rotate /Times-Roman findfont 100 scalefont setfont 0 0 moveto 0.9 setgray (#1) show grestore}def end}}  \\newcounter{onefig} \\newcounter{fignumber} \\newcount\\nocaptions \\newcount\\nofigures \\newcount\\figwidth \\newcount\\viewgraphs \\def\\figsfile{} \\def\\figspath{} \\def\\paper{} \\def\\format{} \\def\\figlabel{} \\long\\def\\nextfig#1{\\setcounter{figure}{\\value{fignumber}} \\addtocounter{fignumber}{1} \\ifnum \\viewgraphs=1 \\newpage \\pagestyle{empty} \\fi \\ifnum\\value{onefig}=0 #1 \\fi \\ifnum\\value{onefig}=\\value{fignumber} #1 \\fi} \\def\\figwidths#1#2{\\ifnum \\nocaptions=1 #2mm \\else #1mm \\fi} \\def\\paper#1{}  % \\long\\def\\plotfig#1#2{\\ifnum \\nofigures=1 \\else #2 \\fi} \\long\\def\\captiontext#1{\\ifnum \\nofigures=1 \\raggedright \\fi \\ifnum \\nocaptions=1 \\paper \\ifnum \\viewgraphs=0 \\newline  \\mbox{}\\hrulefill\\mbox{} \\newline \\newline label:~\\{\\figlabel\\} \\fi \\else \\ifnum \\nofigures=0 \\fi #1 \\fi}  \\newcount\\panelwidth \\newcount\\panelheight \\newcount\\bxmin \\newcount\\bymin \\newcount\\bxmax \\newcount\\bymax \\newcount\\tbxmin \\newcount\\tbymin \\newcount\\tpanelwidth \\newcount\\tpanelheight \\newcount\\tpdif \\panelwidth=70 \\panelheight=70  % \\def\\panelsize #1,#2;{\\panelwidth=#1 \\panelheight=#2} \\def\\setbb #1,#2;#3,#4;#5,#6;{% \\tbxmin=#1 \\tbymin=#2    % \\bxmin=#3 \\bymin=#4      % \\bxmax=#5 \\bymax=#6}     % \\def\\barepanel #1{% \\ifnum\\panelheight=0 \\tpdif=\\bymax \\advance\\tpdif by -\\bymin \\multiply \\tpdif by \\panelwidth \\tpanelheight=\\tpdif \\tpdif=\\bxmax \\advance\\tpdif by -\\bxmin \\divide \\tpanelheight by \\tpdif \\else \\tpanelheight=\\panelheight \\fi \\epsfig{file=#1,% bbllx=\\bxmin bp,bblly=\\bymin bp,bburx=\\bxmax bp,bbury=\\bymax bp,clip=,% width=\\panelwidth mm,height=\\tpanelheight mm}} \\def\\labelypanel #1{% \\ifnum\\panelheight=0 \\tpdif=\\bymax \\advance\\tpdif by -\\bymin \\multiply \\tpdif by \\panelwidth \\tpanelheight=\\tpdif \\tpdif=\\bxmax \\advance\\tpdif by -\\bxmin \\divide \\tpanelheight by \\tpdif \\else \\tpanelheight=\\panelheight \\fi \\tpdif=\\bxmax \\advance\\tpdif by -\\tbxmin \\tpanelwidth=\\panelwidth \\multiply \\tpanelwidth by \\tpdif \\tpdif=\\bxmax \\advance\\tpdif by -\\bxmin \\divide \\tpanelwidth by \\tpdif \\epsfig{file=#1,% bbllx=\\tbxmin bp,bblly=\\bymin bp,bburx=\\bxmax bp,bbury=\\bymax bp,% clip=,width=\\tpanelwidth mm,height=\\tpanelheight mm}} \\def\\labelxpanel #1{% \\ifnum\\panelheight=0 \\tpdif=\\bymax \\advance\\tpdif by -\\bymin \\multiply \\tpdif by \\panelwidth \\tpanelheight=\\tpdif \\tpdif=\\bxmax \\advance\\tpdif by -\\bxmin \\divide \\tpanelheight by \\tpdif \\else \\tpanelheight=\\panelheight \\fi \\tpdif=\\bymax \\advance\\tpdif by -\\tbymin \\multiply \\tpanelheight by \\tpdif \\tpdif=\\bymax \\advance\\tpdif by -\\bymin \\divide \\tpanelheight by \\tpdif \\epsfig{file=#1,% bbllx=\\bxmin bp,bblly=\\tbymin bp,bburx=\\bxmax bp,bbury=\\bymax bp,% clip=,width=\\panelwidth mm,height=\\tpanelheight mm}} \\def\\labelxypanel #1{% \\ifnum\\panelheight=0 \\tpdif=\\bymax \\advance\\tpdif by -\\bymin \\multiply \\tpdif by \\panelwidth \\tpanelheight=\\tpdif \\tpdif=\\bxmax \\advance\\tpdif by -\\bxmin \\divide \\tpanelheight by \\tpdif \\else \\tpanelheight=\\panelheight \\fi \\tpdif=\\bxmax \\advance\\tpdif by -\\tbxmin \\tpanelwidth=\\panelwidth \\multiply \\tpanelwidth by \\tpdif \\tpdif=\\bxmax \\advance\\tpdif by -\\bxmin \\divide \\tpanelwidth by \\tpdif \\tpdif=\\bymax \\advance\\tpdif by -\\tbymin \\multiply \\tpanelheight by \\tpdif \\tpdif=\\bymax \\advance\\tpdif by -\\bymin \\divide \\tpanelheight by \\tpdif \\epsfig{file=#1,% bbllx=\\tbxmin bp,bblly=\\tbymin bp,bburx=\\bxmax bp,bbury=\\bymax bp,% clip=,width=\\tpanelwidth mm,height=\\tpanelheight mm}}  \\newcommand{\\panellabel}[3]{\\makebox[0pt][r]{\\raisebox{#2} {#3} \\hspace{#1}}}  \\def\\topfraction{0.99}                      % \\def\\dbltopfraction{0.99} \\def\\textfraction{0.01} \\def\\floatpagefraction{0.99}  \\def\\CC{\\par \\vspace*{-2ex} \\footnotesize \\baselineskip=8pt \\begin{verbatim}} \\def\\EE{\\par \\vspace*{-2ex} \\normalsize}  \\long\\def\\startignore #1\\stopignore{}   %   \\newenvironment{task}{% \\begin{list}{$\\bullet$}{\\itemsep=1ex \\parindent=3ex \\parskip=1ex \\topsep=0ex}} {\\end{list}}  \\def\\koprr{\\list{}{ \\itemindent=0pt \\listparindent=0pt \\itemsep=0pt \\leftmargin=5ex}\\item[]}                     % \\let\\endkoprr=\\endlist  \\def\\kolom{\\list{}{\\rightmargin=5cm      % \\topsep=0ex \\itemsep=0cm \\parsep=0ex \\parskip=0ex \\partopsep=0ex \\listparindent=3ex \\itemindent=0pt \\leftmargin=0pt} \\item[] \\vspace*{-0ex}}                        % \\let\\endkolom=\\endlist  \\renewcommand{\\labelitemi}{{\\bf --}}                % \\def\\setlistparams{ \\topsep=0.7ex                 % \\itemsep=0.7ex                % \\leftmargini=3ex}             % \\setlistparams                  %  \\newcounter{alistindex}  \\newcounter{alistvalue}     % \\newenvironment{alist}{% \\begin{list}{\\alph{alistindex})}{\\usecounter{alistindex}} \\itemsep=1.5ex \\parskip=1ex}{\\end{list}}  \\def\\spacing #1{\\small \\renewcommand{\\baselinestretch}{#1} \\normalsize \\setlistparams}             %  \\newenvironment{itemrr}{\\vspace*{-\\parskip} \\leftmargini=3ex \\topsep=0.6ex \\begin{itemize} \\itemsep=0.5ex \\parindent=0ex \\parsep=0ex}{ \\end{itemize} \\vspace*{-\\parskip}}  \\newenvironment{conditem}{ \\leftmargini=6ex \\topsep=0.2ex \\begin{itemize} \\itemsep=0ex \\parindent=3ex \\parsep=0ex}{ \\end{itemize}}  \\newenvironment{enumerr}{ \\leftmargini=3ex \\topsep=1.5ex \\begin{enumerate} \\itemsep=0.7ex \\parindent=3ex \\parsep=0ex}{ \\end{enumerate}}  \\newenvironment{condenum}{ \\leftmargini=3ex \\topsep=0.5ex \\begin{enumerate} \\itemsep=0ex \\parindent=3ex \\parsep=0ex}{ \\end{enumerate}}  \\newcounter{romenumnr} \\newenvironment{romanenum}{ \\leftmargini=5ex \\topsep=1.5ex \\begin{list}{(\\roman{romenumnr}) --}{\\usecounter{romenumnr} \\itemsep=1ex \\listparindent=3ex \\parsep=0ex}}{ \\end{list}}  \\newcommand{\\posterpage}[2]{{\\Large \\sf \\newpage \\setlength{\\minipagewidth}{\\textwidth} \\addtolength{\\minipagewidth}{-2cm} \\fboxrule=0.3ex \\fboxsep=0ex \\def\\shadowparams{\\topsep=-2.5ex \\itemsep=-1.5ex \\leftmargini=5ex} \\renewcommand{\\labelitemi}{$\\bullet$} \\parskip=3ex \\shadowparams \\savebox{\\boxcontent}[\\textwidth]{ \\begin{minipage}[t]{\\minipagewidth} \\parskip=3ex \\vspace*{1ex} \\begin{center} \\LARGE \\bf #1 \\end{center} \\vspace{2ex} #2 \\vspace*{3ex} \\mbox{ } \\end{minipage}} \\hspace*{\\textwidth} \\hspace*{-0.8ex} \\rule{2ex}{\\dp\\boxcontent} \\par \\vspace*{-\\dp\\boxcontent} \\vspace*{-7.2ex} \\framebox[\\textwidth]{\\usebox{\\boxcontent}} \\par \\vspace*{-3.0ex} \\hspace*{1.7ex} \\rule{\\textwidth}{2ex} \\setlistparams \\renewcommand{\\labelitemi}{{\\bf --}}}}  \\newcommand{\\vwgraph}[3]{\\newpage \\begin{center} \\shadowframe{#1}{ \\begin{center} \\vspace*{2ex} {\\Large \\bf #2}\\\\[4ex] \\end{center} {\\Large \\sf #3}}   % \\end{center}} \\newenvironment{itemvg}{ \\leftmargini=3ex \\topsep=0.5ex \\begin{itemize} \\vspace{-1ex} \\itemsep=0.5ex \\parindent=3ex \\parsep=0ex}{ \\end{itemize}}  \\long\\def\\lectvw#1#2{\\newpage \\vspace*{-15mm} \\begin{center} \\begin{center} {\\LARGE \\sc \\underline{#1}}\\\\[4ex] \\end{center} {\\Large \\sf #2}   % \\end{center}}  \\newlength{\\minipagewidth} \\newcommand{\\shadowframe}[2]{ \\setlength{\\minipagewidth}{#1} \\addtolength{\\minipagewidth}{-1cm} \\fboxrule=0.5ex \\fboxsep=0ex \\def\\shadowparams{\\topsep=-2.5ex \\itemsep=-1.5ex \\leftmargini=5ex} \\renewcommand{\\labelitemi}{$\\bullet$} \\parskip=3ex \\shadowparams \\savebox{\\boxcontent}[#1]{ \\begin{minipage}[t]{\\minipagewidth} \\vspace*{1.5ex} #2 \\vspace*{0.5ex} \\mbox{ } \\end{minipage}} \\hspace*{#1} \\hspace*{1.0ex} \\rule{2ex}{\\dp\\boxcontent} \\par \\vspace*{-\\dp\\boxcontent} \\vspace*{-7.9ex} \\framebox[#1]{\\usebox{\\boxcontent}} \\par \\vspace*{-3.0ex} \\hspace*{4.5ex} \\rule{#1}{2ex} \\setlistparams \\renewcommand{\\labelitemi}{{\\bf --}}}  \\newsavebox{\\boxcontent} \\newcommand{\\ovalhead}[1]{ \\unitlength=1cm \\sbox{\\boxcontent}{\\mbox{~~{#1}~~}} \\begin{center} \\ifdim\\wd\\boxcontent>6ex \\ifdim\\wd\\boxcontent<8cm \\begin{picture}(8,3) \\thicklines \\put(4.0,0.8){\\oval(8,1.6)} \\put(0.0,0.7){\\parbox{8cm}{ \\begin{center} \\usebox{\\boxcontent} \\end{center}}} \\end{picture} \\else \\ifdim\\wd\\boxcontent<12cm \\begin{picture}(12,3) \\thicklines \\put(6.0,0.8){\\oval(12,1.6)} \\put(0.0,0.7){\\parbox{12cm}{ \\begin{center} \\usebox{\\boxcontent} \\end{center}}} \\end{picture} \\else \\begin{picture}(16,3) \\thicklines \\put(8.0,0.8){\\oval(16,1.6)} \\put(0.0,0.7){\\parbox{16cm}{ \\begin{center} \\usebox{\\boxcontent} \\end{center}}} \\end{picture} \\fi \\fi \\fi \\end{center}}  \\newcommand{\\colloqframe}[2]{ \\setlength{\\minipagewidth}{#1} \\addtolength{\\minipagewidth}{-1cm} \\fboxrule=0.5ex \\fboxsep=0ex \\parskip=3ex \\savebox{\\boxcontent}[#1]{ \\begin{minipage}[t]{\\minipagewidth} \\vspace*{1.5ex} #2 \\vspace*{0.5ex} \\mbox{ } \\end{minipage}} \\hspace*{#1} \\hspace*{-0.8ex} \\rule{2ex}{\\dp\\boxcontent} \\par \\vspace*{-\\dp\\boxcontent} \\vspace*{-7.9ex} \\framebox[#1]{\\usebox{\\boxcontent}} \\par \\vspace*{-3.0ex} \\hspace*{1.7ex} \\rule{#1}{2ex} \\renewcommand{\\labelitemi}{{\\bf --}}}  \\newcommand{\\soloval}[1]{\\fboxrule=0.5mm \\unitlength=1cm \\hspace*{5mm} \\sol{20mm}\\\\[-20mm] \\begin{picture}(15,3) \\thicklines \\put(7.5,2){\\oval(15.60,3.60)} \\put(7.5,2){\\oval(15.55,3.55)} \\put(7.5,2){\\oval(15.50,3.50)} \\put(7.5,2){\\oval(15.45,3.45)} \\put(7.5,2){\\oval(15.40,3.40)} \\put(7.5,2){\\oval(15.35,3.35)} \\put(7.5,2){\\oval(15.00,3.00)} \\put(2.0,1.9){\\parbox{13cm}{#1}} \\end{picture}}  \\setcounter{secnumdepth}{3} \\setcounter{tocdepth}{3}  \\def\\chapterrr#1{\\chapter{#1} \\label{chap:#1} \\vskip-\\parskip} \\def{} {We aim to determine the fundamental parameters of a sample of B stars with apparent visual magnitudes below 8 in the field-of-view of the CoRoT space mission, from high-resolution spectroscopy. } {We developed an automatic procedure for the spectroscopic analysis of B-type stars with winds, based on an extensive grid of FASTWIND model atmospheres.  We use the equivalent widths and/or the line profile shapes of continuum normalized hydrogen, helium and silicon line profiles to determine the fundamental properties of these stars in an automated way. } {After thorough tests, both on synthetic datasets and on very high-quality, high-resolution spectra of B stars for which we already had accurate values of their physical properties from alternative analyses, we applied our method to 66 B-type stars contained in the ground-based archive of the CoRoT space mission. We discuss the statistical properties of the sample and compare them with those predicted by evolutionary models of B stars.} {Our spectroscopic results provide a valuable starting point for any future seismic modelling of the stars, should they be observed by CoRoT.} ",
        "introduction": "The detailed spectroscopic analysis of B-type stars has for a long time been restricted to a limited number of targets. Reasons for this are the a priori need for a realistic atmosphere model, the lack of large samples with high-quality spectra, and the long-winded process of line profile fitting, as the multitude of photospheric and wind parameters requires a large parameter space to be explored.  The advent of high-resolution, high signal-to-noise spectroscopy in the nineties led to a renewed interest of the scientific community in spectroscopic research, and in particular in the relatively poorly understood massive stars. The establishment of continuously better instrumentation and the improvement in quality of the obtained spectroscopic data triggered a series of studies, which led to a rapid increase in our knowledge of massive stars. In this respect, it is not surprising to note that this is exactly the period where several groups started to upgrade their atmosphere prediction code for such stars, see, e.g., CMFGEN - \\citet{Hillier1998}, PHOENIX - \\citet{Hauschildt1999}, WM-Basic - \\citet{Pauldrach2001}, POWR - \\citet{Grafener2002} and FASTWIND - \\citet{Santolaya1997}, \\citet{Puls2005}. Initially, major attention was devoted to the establishment of a realistic atmosphere model (improvement of atomic data, inclusion of line blanketing and clumping), rather than analyzing large samples of stars.  Simultaneously with improvements in the atmosphere predictions, also the number of available high-quality data increased rapidly, mainly thanks to the advent of multi-object spectroscopy. At the time of writing, the largest survey is the VLT-FLAMES Survey of Massive Stars \\citep{Evans2005}, containing over 600 Galactic, SMC and LMC B-type spectra\\footnote{additionally, roughly 90 O-stars have been observed.} (in 7 different clusters), gathered over more than 100 hours of VLT time. The survey not only allowed to derive the stellar parameters and rotational velocities for hundreds of stars \\citep{Dufton2006, Hunter2008}, but also to study the evolution of surface N abundances and the effective temperature scales in the Galaxy and Magellanic Clouds \\citep{Trundle2007, Hunter2007, Hunter2008a}.  In preparation of the CoRoT space mission, and almost contemporary with the FLAMES setup, another large database was constructed: GAUDI (Ground-based Asteroseismology Uniform Database Interface, \\citealt{Solano2005}). It gathers ground-based observations of more than 1500 objects, including high-resolution spectra of about 250 massive B-type stars, with the goal to determine their fundamental parameters as input for seismic modeling (see Section\\,\\ref{GAUDI}).  The availability of such large samples of B-type stars brings within reach different types of studies, e.g., they may lead to a significant improvement in the fundamental parameter calibration for this temperature range and to a confrontation with and evaluation of stellar evolution models (e.g., \\citealt{Hunter2008}). The drawback of this huge flood of data, however, is, as mentioned before, the large parameter space to be explored, which can be quite time-consuming, if no adequate method is available. It requires a method which is able to derive the complete set of parameters of stars with a wide variety of physical properties in an objective way.  To deal with the large GAUDI dataset, we investigated the possibility of \\textit{automated} spectral line fitting and we opted for a grid-based fitting method: AnalyseBstar. In Section\\,\\ref{method}, we justify our choice for a grid-based method, present its design and discuss several tests which were applied to check the performance of the routine. A more detailed description of our methodology can be found in the (online) appendix. Section\\,\\ref{GAUDI} illustrates the first application of AnalyseBstar to the sample of CoRoT candidate targets in the GAUDI database and Section\\,\\ref{interpretation} deals with the physical interpretation and some statistical properties of the resulting parameters. Section\\,\\ref{summary} summarizes the main results obtained in this paper. ",
        "conclusions": ""
    },
    "0910/0910.2256_arXiv.txt": {
        "abstract": "{ The lowest-mass stars, brown dwarfs and extrasolar planets present challenges and opportunities for understanding dynamics and cloud formation processes in low-temperature atmospheres.  For brown dwarfs, the formation, variation and rapid depletion of photospheric clouds in L- and T-type dwarfs, and spectroscopic evidence for non-equilibrium chemistry associated with vertical mixing, all point to a fundamental role for dynamics in vertical abundance distributions and cloud/grain formation cycles. For exoplanets, azimuthal heat variations and the detection of stratospheric and exospheric layers indicate multi-layered, asymmetric atmospheres that may also be time-variable (particularly for systems with highly elliptical orbits).  Dust and clouds may also play an important role in the thermal energy balance of exoplanets through albedo effects.  For all of these cases, 3D atmosphere models are becoming an increasingly essential tool for understanding spectral and temporal properties.  In this review, I summarize the observational evidence for clouds and dynamics in cool dwarf and hot exoplanetary atmospheres, outstanding problems associated with these processes, and areas where effective synergy can be achieved. ",
        "introduction": "Fifteen years ago, the first examples of brown dwarfs and extrasolar planets were discovered, sources which continue to challenge our understanding of low temperature atmospheres.  Several hundred very low-mass stars and brown dwarfs are now known (the latter distinguished by their lack of core hydrogen fusion), encompassing two newly defined spectral classes, the L dwarfs and the T dwarfs (see review by \\citealt{2005ARA&A..43..195K}).  Roughly 300 extrasolar planets have been found through Doppler shift, transit, microlensing and direct detection techniques. Collectively, these low-temperature sources span a broad range of mass (1~{\\mjup} $\\lesssim$ M $\\lesssim$ 100~{\\mjup}), photospheric temperature (100~K $\\lesssim$ $T$ $\\lesssim$ 2500~K), surface gravity (3 $\\lesssim$ {\\logg} $\\lesssim$ 5.5~cm~s$^{-2}$), age (few Myr to few Gyr), elemental composition, rotation period (hours to days), magnetic activity, degree of external heating and interior structure. Yet all are related by their cool, molecule-rich, dynamic atmospheres.  While direct studies of cool dwarf atmospheres have been feasible since their discovery, it is only recently that techniques to probe exoplanet atmospheres have been realized. The best-constrained exoplanets are those which closely orbit and transit their host stars---so-called ``Hot Jupiters''---allowing reflectance and thermal spectrophotometry during secondary transit (e.g., \\citealt{2005ApJ...626..523C, 2005Natur.434..740D}), and transmission spectrophotometry during primary transit (e.g., \\citealt{2002ApJ...568..377C}). Phase curves for non-transiting systems have also been measured (e.g., \\citealt{2006Sci...314..623H}), and the recent direct detection of planets around the young stars Fomalhaut, $\\beta$ Pictoris and HR~8799 \\citep{2008Sci...322.1345K,2008Sci...322.1348M, 2009A&A...493L..21L} through high contrast imaging techniques have opened the door to direct investigations of exoplanet atmospheres.  Two themes are prominent in the interpretation of observational data for low-temperature stellar, brown dwarf and exoplanetary atmospheres: clouds and dynamics.  Condensate clouds are an important source of opacity and a component in the chemical network; they also modulate the albedos, energy budgets and atmospheric structure of exoplanets under intense irradiation.    Dynamics are responsible for the redistribution of heat and modification of chemical abundances, and are likely fundamental to cloud formation and evolution processes.  In this review, I summarize our current observational evidence for clouds and dynamics in cool dwarfs and hot exoplanets, and note outstanding issues in their role in emergent spectral energy distributions, long-term thermal evolution, albedo, and temporal and azimuthal variability. I conclude with a discussion of opportunities for synergy in future generations of cool dwarf and exoplanet models, with a view toward 3D simulations. ",
        "conclusions": ""
    },
    "0910/0910.1533_arXiv.txt": {
        "abstract": "The stellar populations of galaxies contain a wealth of detailed information. From the youngest, most massive stars, to almost invisible remnants, the history of star formation is encoded in the stars that make up a galaxy. Extracting some, or all, of this information has long been a goal of stellar population studies. This was achieved in the last couple of decades and  it is now a routine task, which forms a crucial ingredient in much of observational galaxy evolution, from our Galaxy out to the most distant systems found.  In many of these domains we are now limited not by sample size, but by systematic uncertainties and this will increasingly be the case in the future.  The aim of this review is to outline the challenges faced by stellar population studies in the coming decade within the context of upcoming observational facilities.  I will highlight the need to better understand the near-IR spectral range and outline the difficulties presented by less well understood phases of stellar evolution such as thermally pulsing AGB stars, horizontal branch stars and the very first stars. The influence of rotation and binarity on stellar population modelling is also briefly discussed. ",
        "introduction": "\\label{sec:introduction}  The luminous output of normal galaxies is ultimately generated by stars in various stages of stellar evolution --- in isolation a trite observation, but this simple fact gives us the opportunity to extract a vast amount of information from observations of galaxies through the modelling of their stellar populations.  While the study of stellar populations in galaxies started with Baade's (1944)\\nocite{1944ApJ...100..137B} identification of two populations of stars in M32 and NGC 205, a rigourous study of the topic only commenced in the late 60's and early 70's \\citep{Tinsley1968,1972A&A....20..361F,1973ApJ...179..427S} with Tinsley's \\textit{Fundamentals of Cosmic Physics} article \\citep{1980FCPh....5..287T} particularly influential.  Since that time, the number of articles discussing stellar populations has risen rapidly so that today about 12\\% of all articles in the major journals mention stellar populations in their abstracts (Figure 1).     The majority of this growth has been made possible through the development of simple models for the evolution of stellar populations that have found widespread use in a wide range of astronomical studies, from stellar clusters in the Milky Way to the most distant galaxies in the Universe. I will repeatedly refer to stellar population models below, and use this term loosely to refer to any model that predict the observational properties of any ensemble of stars. These start from models of stellar evolution (e.g. Bertelli et al 1994; see contribution by Cassisi these proceedings). By applying empirical colour corrections (e.g. Lejeune et al 1997) they can be placed on an observational Hertzsprung-Russell diagram. For a given star formation history and initial mass function (IMF), this can be sampled, either to produce a Monte Carlo realisation of an observed colour-magnitude diagrams (see Tolstoy et al 2009), or by convolution to create integrated properties of a stellar population (e.g. Fioc \\& Rocca-Volmerange 1997; Leitherer et al 1999; Vazdekis 1999; Bruzual \\& Charlot 2003;  Maraston 2005; Kotulla et al 2009). Finally, these models are compared to observations using a number of techniques, in itself an interesting topic, but for this review the details of this step are not crucial.    In this contribution I will begin by briefly reviewing some of the achievements in extra-galactic astronomy made possible by our understanding of stellar populations, before I turn to the challenges for stellar populations studies in the coming decade. Given the format this will not be an exhaustive review, see the other contributions in these proceedings for more in-depth discussion of many of the topics. ",
        "conclusions": "\\label{sec:conclusions}  Rather than a traditional conclusion, I would here like to end by summarising the challenges above into a set of rest wavelength ranges, highlighting the issues and some suggestions as to what might be useful studies to do --- note that most of these suggestions are merely reinforcing already existing studies!  \\textbf{Far-UV}  ---  \\textit{Massive stars --- what are realistic rotation and binary parameters}. \\textbf{To do:}  Observations of lensed Lyman-break galaxies, COS observations of nearby stars and star-forming regions,  more in-depth studies of stellar mass-loss  \\textbf{Near-UV}  ---  \\textit{Binary evolution, mass loss on the RGB}. \\textbf{To do:}  COS and GALEX spectroscopy of UV-upturn galaxies, inclusion of horizontal branch uncertainties in models. Combination of optical, near-UV and X-ray data in analysis.   \\textbf{Optical} --- \\textit{High-precision predictions. Non-solar abundance variations and impact on spectra,  emission lines}. \\textbf{To do:} Careful and extensive testing of theoretical spectra, comparison to detailed spectroscopic analysis of nearby systems: what abundances can be reliably extracted from medium resolution spectra?   \\textbf{Near-IR}  ---  \\textit{Evolution of and spectra of luminous, cold stars. TP-AGB stars, RGs and RSGs, spectral features in the near-IR}. \\textbf{To do:} Further observations of resolved stellar populations in the near-IR with and without adaptive optics with careful comparison to results including optical data. Calibration spectral features in the near-IR. Obtain ACS optical data to back up future studies in near-IR with ELTs. In depth analysis of uncertainties in population synthesis models and careful comparisons to data in the nearby and distant Universe.      \\begin{multicols}{2}"
    },
    "0910/0910.5917_arXiv.txt": {
        "abstract": "The presence of convective motions in the atmospheres of metal-poor halo stars leads to systematic asymmetries of the emergent spectral line profiles. Since such line asymmetries are very small, they can be safely ignored for standard spectroscopic abundance analysis. However, when it comes to the determination of the $^6$Li/$^7$Li isotopic ratio, $q$(Li)=$n$($^6$Li)/n($^7$Li), the intrinsic asymmetry of the $^7$Li line must be taken into account, because its signature is essentially indistinguishable from the presence of a weak $^6$Li blend in the red wing of the $^7$Li line. In this contribution we quantity the error of the inferred $^6$Li/$^7$Li isotopic ratio that arises if the convective line asymmetry is ignored in the fitting of the $\\lambda\\,6707$~\\AA\\ lithium blend. Our conclusion is that $^6$Li/$^7$Li ratios derived by \\cite[Asplund \\etal\\ (2006)]{A2006}, using symmetric line profiles, must be reduced by typically $\\Delta q$(Li) $\\approx 0.015$. This diminishes the number of certain $^6$Li detections from 9 to 4 stars or less, casting some doubt on the existence of a $^6$Li plateau. ",
        "introduction": "The spectroscopic signature of the presence of $^6$Li in the atmospheres of metal-poor halo stars is a subtle extra depression in the red wing of the $^7$Li doublet, which can only be detected in spectra of the highest quality. Based on high-resolution, high signal-to-noise VLT/UVES spectra of 24 bright metal-poor stars, \\cite[Asplund \\etal\\ (2006)]{A2006} report the detection of $^6$Li in nine of these objects. The average $^6$Li/$^7$Li isotopic ratio in the nine stars in which $^6$Li has been detected is $q$(Li) $\\approx 0.04$ and is very similar in each of these stars, defining a $^6$Li plateau at approximately $\\log n(^6$Li$) = 0.85$ (on the scale $\\log n($H$) = 12$). A convincing theoretical explanation of this new $^6$Li plateau turned out to be problematic. Even when the depletion of the $^6$Li isotope during the pre-main-sequence phase would be ignored, the high abundances of $^6$Li at the lowest metallicities cannot be explained by current models of galactic cosmic-ray production (for a concise review see e.g.\\ \\cite[Christlieb 2008]{C2008}, and references therein).  A possible solution of the so-called `second Lithium problem' was suggested by \\cite[Cayrel \\etal\\ (2007)]{C2007}, who point out that the intrinsic line asymmetry caused by convection in the photospheres of cool stars is almost indistinguishable from the asymmetry produced by a weak $^6$Li blend on a presumed symmetric $^7$Li profile. As a consequence, the derived $^6$Li abundance should be significantly reduced when the intrinsic line asymmetry in properly taken into account. Using 3D non-LTE line formation calculations based on 3D hydrodynamical model atmospheres computed with the CO$^5$BOLD code (\\cite[Freytag \\etal\\ 2002]{F2002}, \\cite[Wedemeyer \\etal\\ 2004]{W2004}, see also {\\tt http://www.astro.uu.se/$\\sim$bf/co5bold\\_main.html}), we quantify the theoretical effect of the convection-induced line asymmetry on the resulting $^6$Li abundance as a function of effective temperature, gravity, and metallicity, for a parameter range that covers the stars of the \\cite[Asplund \\etal\\ (2006)]{A2006} sample. ",
        "conclusions": "The present study indicates that only $2$ or at most $4$ out of the $24$ stars of the \\cite{A2006} sample remain significant $^6$Li detections when subjected to a 3D~non-LTE analysis, suggesting that the presence of $^6$Li in the atmospheres of galactic halo stars is rather the exception than the rule. This would imply that it is no longer necessary to look for a global mechanism accounting for a $^6$Li enrichment of the galactic halo, but that it is sufficient to explain only a few exceptional cases, which is probably much easier.\\\\[-5mm]"
    },
    "0910/0910.2706_arXiv.txt": {
        "abstract": "Many astrophysical sources, especially compact accreting sources, show strong, random brightness fluctuations with broad power spectra in addition to periodic or quasi-periodic oscillations (QPOs) that have narrower spectra. The random nature of the dominant source of variance greatly complicates the process of searching for possible weak periodic signals. We have addressed this problem using the tools of Bayesian statistics; in particular using Markov chain Monte Carlo techniques to approximate the posterior distribution of model parameters, and posterior predictive model checking to assess model fits and search for periodogram outliers that may represent periodic signals. The methods developed are applied to two example datasets, both long \\xmm\\ observations of highly variable Seyfert 1 galaxies: RE J$1034+396$ and Mrk $766$. In both cases a bend (or break) in the power spectrum is evident. In the case of RE J$1034+396$ the previously reported QPO is found but with somewhat weaker statistical significance than reported in previous analyses. The difference is due partly to the improved continuum modelling, better treatment of nuisance parameters, and partly to different data selection methods. ",
        "introduction": "\\label{sect:intro}  \\defcitealias{Vaughan05}{V05}  A perennial problem in observational astrophysics is detecting periodic or almost-periodic signals in noisy time series. The standard analysis tool is the periodogram \\citep[see e.g. ][]{Jenkins69, Priestley81, Press92, Bloomfield00, Chatfield03}, and the problem of period detection amounts to assessing whether or not some particular peak in the periodogram is due to a periodic component or a random fluctuation in the noise spectrum \\citep[see][]{Fisher29, Priestley81, Leahy83, vanderklis89, Percival93, Bloomfield00}.  If the time series is the sum of a random (stochastic) component and a periodic one we may write $y(t) = y_R(t) + y_P(t)$ and, due to the independence of $y_R(t)$ and $y_P(t)$, the power spectrum of $y(t)$ is the sum of the two power spectra of the random and stochastic processes: $S_Y(f) = S_R(f) + S_P(f)$. This is a \\emph{mixed} spectrum \\citep[][section 4.4]{Percival93} formed from the sum of $S_P(f)$, which comprises only narrow features, and $S_R(f)$, which is a continuous, broad spectral function. Likewise, we may consider an evenly sampled, finite time series $y(t_i)$ ($i=1,2,\\ldots,N$) as the sum of two finite time series: one is a realisation of the periodic process, the other a random realisation of the stochastic process. We may compute the periodogram (which is an estimator of the true power spectrum) from the squared modulus of the Discrete Fourier Transform (DFT) of the time series, and, as with the power spectra, the periodograms of the two processes add linearly: $I(f_j) = I_R(f_j) + I_P(f_j)$. The periodogram of the periodic time series will contain only narrow ``lines'' with all the power concentrated in only a few frequencies, whereas the periodogram of the stochastic time series will show power spread over many frequencies. Unfortunately the periodogram of stochastic processes fluctuates wildly around the true power spectrum, making it difficult to distinguish random fluctuations in the noise spectrum from truly spectral periodic components. See \\cite{vanderklis89} for a thorough review of these issues in the context of X-ray astronomy.  Particular attention has been given to the special case that the spectrum of the stochastic process is flat (a \\emph{white noise} spectrum $S(f) = const$), which is the case when the time series data $y_R(t_i)$ are independently and identically distributed (IID) random variables. Reasonably well-established statistical procedures have been developed to help identify spurious spectral peaks and reduce the chance of false detections \\citep[e.g.][]{Fisher29, Priestley81, Leahy83, vanderklis89, Percival93}. In contrast there is no comparably well-established procedure in the general case that the spectrum of the stochastic process is not flat.  In a previous paper, \\cite{Vaughan05} (henceforth \\citetalias{Vaughan05}), we proposed what is essentially a generalisation of Fisher's method to the case where the noise spectrum is a power law: $S_R(f) = \\beta f^{-\\alpha}$ (where $\\alpha$ and $\\beta$ are the power law index and normalisation parameters). Processes with power spectra that show a power law dependence on frequency with $\\alpha > 0$ (i.e. increasing power to lower frequencies) are called \\emph{red noise} and are extremely common in astronomy and elsewhere \\citep[see][]{Press78}. In this paper we expand upon the ideas in \\citetalias{Vaughan05} and, in particular, address the problem from a Bayesian perspective that allows further generalisation of the spectral model of the noise.  The rest of this paper is organised as follows. In section~\\ref{sect:bayes} we introduce some of the basic concepts of the Bayesian approach to statistical inference; readers familiar with this topic may prefer to skip this section. Section~\\ref{sect:stat} gives a brief overview of classical significance testing using $p$-values (tail area probabilities) and test statistics, and section~\\ref{sect:ppp} discusses the posterior predictive $p$-value, a Bayesian counterpart to the classical $p$-value. Section \\ref{sect:pc} reviews the conventional (classical) approaches to testing for periodogram peaks. Section \\ref{sect:ml} outlines the theory of maximum likelihood estimation from periodogram data, which is developed into the basis of a fully Bayesian analysis in sections~\\ref{sect:ba} and \\ref{sect:ppper}. The Bayesian method is then applied to two real observations if AGN in section \\ref{sect:data}. Section \\ref{sect:disco} discusses the limitations of the method, and alternative approaches to practical data analysis. A few conclusions are given in section \\ref{sect:conc}, and two appendices describe details of the simulations algorithms used in the analysis. ",
        "conclusions": "\\label{sect:conc}  We have presented Bayesian methods for the modelling of periodogram data that can be used for both parameter estimation and model checking, and may be used to test for narrow spectral features embedded in noisy data. The model assessment is performed using simulations of posterior predictive data to calibrate (sensibly chosen) test statistics. This does however leave some arbitrariness in the method, particularly in the choice of test statistic\\footnote{In situations where two competing models can be modelled explicitly the LRT provides a natural choice of statistic.} (and in some situations the choice of what constitutes a simulation of the data). Such issues were always present, if usually ignored, in the standard frequentist tests. The posterior predictive approach has the significant advantage of properly treating nuisance parameters, and provides a clear framework for checking the different aspects of the reasonableness of a model fit. The issue of choosing a test statistic does not arise in more ``purist'' Bayesian methods such as Bayes factors, which concentrate on the posterior distributions and marginal likelihoods, but such methods of model selection carry their own burden in terms of the computational complexity and the difficulty of selecting (and the sensitivity of inferences to) priors on the model parameters. The method presented in this paper, making use of the posterior predictive checking, is an improvement over the currently popular methods that use classical $p$-value; but Bayesian model selection is an area of active research and it is not unreasonable to expect that new, powerful and practical computational tools will be developed or adapted to help bridge the gap between the pragmatic and the purist Bayesian approaches.  The routines used to perform the analysis of the real data presented in section \\ref{sect:data} will be made available as an {\\tt R}\\footnote{{\\tt R} is a powerful, open-source computing environment for data analysis and statistics that may be downloaded for free from {\\tt http://www.r-project.org/} \\citep{r, r2}.} script from the author upon request."
    },
    "0910/0910.0159_arXiv.txt": {
        "abstract": "We investigate the case of [{\\sc Cii}] 158$\\mu$m observations for SPICA/SAFARI using a three-dimensional magnetohydrodynamical (MHD) simulation of the diffuse interstellar medium (ISM) and the Meudon PDR code. The MHD simulation consists of two converging flows of warm gas ($10^4$ K) within a cubic box 50 pc in length. The interplay of thermal instability, magnetic field and self-gravity leads to the formation of cold, dense clumps within a warm, turbulent interclump medium. We sample several clumps along a line of sight through the simulated cube and use them as input density profiles in the Meudon PDR code. This allows us to derive intensity predictions for the [{\\sc Cii}] 158$\\mu$m line and provide time estimates for the mapping of a given sky area. ",
        "introduction": "The fine structure line of ionised carbon [{\\sc Cii}] at 157.7~$\\mu$m is one of the most prominent far infrared (FIR) line. It is widespread in the Milky Way as revealed by the COBE-FIRAS observations, and detected in the diffuse ISM \\citep{levrierf:I02}, photo-dissociation regions as well as local and distant galaxies \\citep[and references therein]{levrierf:M09}. [{\\sc Cii}] is one of the main cooling lines of the neutral ISM, in all environments where carbon is mostly ionised, that is where UV-photons with $h\\nu \\geq 11.3~\\mathrm{eV}$ are available. Unfortunately, the present information on the spatial structure of the [{\\sc Cii}] emission is very limited, because of the small telescope size of previous far infrared missions. While Herschel and SOFIA will improve the situation for bright regions, their sensitivity will be limited for the extended emission from either the diffuse interstellar medium, or the outskirts of molecular clouds. The structure of the [{\\sc Cii}] emission is expected to reveal the underlying structure of the matter, and provide information on the mechanisms responsible for the fragmentation. In this paper, we explore the perspectives offered by SPICA for mapping the interstellar medium. ",
        "conclusions": "This work demonstrates the ability of SPICA/SAFARI to map large areas of the sky in the [{\\sc Cii}] line in a much shorter time than what would be possible with Herschel, making it clear that the extraordinary sensitivity of the proposed mission is a definitive asset regarding this type of project.  On another level, the work presented here is part of an ongoing ASTRONET project dubbed STAR FORMAT, which is a German-French collaboration whose two main objectives are :  1/ Producing databases to publicize simulation results from MHD and PDR codes;  2/ Developing MHD, PDR and radiative transfer codes in an interoperable fashion, to allow for a much improved treatment of the physics of the ISM, especially regarding the formation of molecular clouds and dense cores.  With respect to this project, our work shows that a truly three-dimensional PDR code is a must, to adequately deal with the complex geometries of the structures formed in the MHD simulations. The development of parallel computing schemes is also quite inevitable, and it should be noted that the PDR code can already be run on a grid. It is hoped that the present paper can be used as ground work for these developments."
    },
    "0910/0910.0968_arXiv.txt": {
        "abstract": "The geoeffective magnetic cloud (MC) of 20 November 2003, has been associated to the 18 November 2003, solar active events in previous studies. In some of these, it was estimated that the magnetic helicity carried by the MC had a positive sign, as well as its solar source, active region (AR) NOAA 10501. In this paper we show that the large-scale magnetic field of AR 10501 had a negative helicity sign. Since coronal mass ejections (CMEs) are one of the means by which the Sun ejects magnetic helicity excess into the interplanetary space, the signs of magnetic helicity in the AR and MC should agree. Therefore, this finding contradicts what is expected from magnetic helicity conservation. However, using for the first time correct helicity density maps to determine the spatial distribution of magnetic helicity injection, we show the existence of a localized flux of positive helicity in the southern part of AR 10501. We conclude that positive helicity was ejected from this portion of the AR leading to the observed positive helicity MC. ",
        "introduction": "\\label{S-Intro}  Magnetic helicity globally quantifies the signed amount of twist, writhe, and shear of the magnetic field in a given volume (see the review by \\inlinecite{Demoulin07} and references therein). Magnetic helicity plays an important role in magnetohydrodynamics (MHD) because it is one of the few global quantities which are conserved, even in resistive MHD on  time scales shorter than the global diffusion time scale (\\opencite{Berger84}).  Coronal mass ejections (CMEs) are expulsions of mass and magnetic field from the Sun. \\inlinecite{Rust94} and \\inlinecite{Low96} pointed out that one of the most important roles of CMEs is to carry away magnetic helicity from the Sun. Otherwise, because helicity dissipates very slowly and helicities of opposite sign are globally injected through the photosphere in each solar hemisphere without change of sign during consecutive cycles (\\opencite{Berger00}), helicity would accumulate continuously. A fraction of CMEs can be observed {\\it in situ} as magnetic clouds (MCs). An MC is characterized by lower proton temperature and higher magnetic field strength than the surrounding solar wind.  Typically, the magnetic field vector shows a smooth and significant rotation across the cloud (\\opencite{Burlaga81}; \\opencite{Klein82}) indicating a helical (flux rope) magnetic structure, which clearly has non-zero helicity. Therefore, a measurable prediction is that the interplanetary MC must carry the same amount of helicity that was ejected from the solar source region. In particular, the signs of the magnetic helicity of the MC and of the solar region from which it originates should agree.  As a first approach, the magnetic helicity sign of some structures in the solar atmosphere can be inferred from certain observed morphological features (see the review by \\inlinecite{Demoulin09} and references therein). These features include {\\it sunspot whorls} (handedness of the spiral patterns of chromospheric fibrils), {\\it filament barbs} (direction of the barbs relative to the orientation of the magnetic field), {\\it flare-ribbons} (forward/reverse `J-shape' observed in H$\\alpha$ and UV wavelengths), {\\it sigmoids} (normal or reverse `S-shaped' loops in soft X-ray observations), {\\it magnetic tongues} (angle formed by the magnetic inversion line relative to the AR axis in emerging ARs), {\\it coronal loops} (orientation of loops in EUV observations relative to the magnetic inversion line), {\\it vector magnetic field} (direction of sheared fields). These observational features can be used to qualitatively compare the magnetic helicity sign of an MC with that of its solar source once identified ({\\it e.g.} \\opencite{Subramanian01}; \\opencite{Schmieder05}; \\opencite{Ali07}). In a similar way, {\\it i.e.} directly from observations, the helicity sign of MCs can be estimated from the measured rotation of the vector magnetic field (see examples in Bothmer and Schwenn (1994, 1998)).  Several studies have found that the magnetic helicity sign of the MC and that inferred from the morphological features of its source AR match ({\\it e.g.} \\opencite{RK94}; Bothmer and Schwenn, 1994, 1998; \\opencite{Marubashi97}; \\opencite{Ruzmaikin03}; \\opencite{Rust05}). Other studies have determined the helicity sign of the CME source region modeling the coronal magnetic field ({\\it e.g.} \\opencite{Yurchyshyn01}, \\citeyear{Yurchyshyn06}), in these cases also the source region and cloud helicity signs were in agreement. However, the comparison does not yield a complete agreement, since a few MCs seem to present a different helicity sign from that of their solar source \\cite{Leamon04}.  Quantitative comparisons of the helicity involved in the solar ejection and that of the associated MC have recently been possible. In the interplanetary medium, these quantitative comparisons require either the modeling of {\\it in situ} magnetic field observations (see the reviews by \\inlinecite{Dasso05} and \\inlinecite{Nakwacki08}) or, in cases when the impact parameter is small and considering a local cylindrical geometry, the MC helicity can be directly quantified from the data \\cite{Dasso06}. In the solar atmosphere, at least two different methods, giving consistent results (\\opencite{Lim07}), allow the estimation of the amount of ejected helicity. One way is to compute the helicity variation before and after the ejection of the solar source region using a coronal field model (\\opencite{Green02}; \\opencite{Demoulin02}; \\opencite{Mandrini05}; \\opencite{Regnier05}; \\opencite{Luoni05}). This method requires the knowledge of the magnetic field in the entire volume. Another way is to measure the helicity injection, based on a time series of photospheric field observations. This was initiated by \\inlinecite{Chae01b} and \\inlinecite{Chae01}, and was subsequently applied to the study of CMEs by \\inlinecite{Nindos02} and \\inlinecite{Nindos03}. \\inlinecite{Pariat05} showed that, although the methods used in previous works could correctly estimate the total injected flux of helicity, they incorrectly determined the localized injected flux and proposed an alternative to properly map the helicity flux injection. We shall use this corrected method to compute the magnetic helicity injection in AR 10501.  The large MC observed on 20 November 2003, gave place to the largest geomagnetic storm of solar cycle 23 \\cite{Gopalswamy05}. This MC was associated to the active solar events that occurred in AR 10501 on 18 November 2003, by several authors (\\opencite{Gopalswamy05}; \\opencite{Yurchyshyn05}; \\opencite{Mostl08}). The association discussed by \\inlinecite{Gopalswamy05} is mainly based on the timing between the filament ejections in AR 10501, the appearance of CMEs in the Large Angle and Spectroscopic Coronagraph (LASCO) white light images (its height-time plots), and the arrival of the MC to the Advanced Composition Explorer (ACE) spacecraft and its velocity. \\inlinecite{Yurchyshyn05} also estimated the AR helicity sign modeling the AR coronal magnetic field and the MC {\\it in situ} data; these latter authors showed that the helicity of the MC and AR were both positive. However, \\inlinecite{Mostl08} discussed that while the MC helicity sign seems well determined, the handedness (or helicity sign) of the very extended filament, lying along different portions of the inversion line within and in the surroundings of the AR, is ambiguous.  In this paper, we revisit the evolution of the activity in AR 10501 along 18 November 2003, in Section~\\ref{S-obs}. Then, we discuss the characteristics of the positive-helicity-carrier MC observed by ACE on 20 November 2003, and its association with the CMEs originating from the AR (Section~\\ref{Ss-MC}). To verify this association, we analyze carefully the morphological features of the region; in particular, the different segments of filament material lying along the magnetic inversion line that are observed as forming a single and very extended filament. We find that all these observational features, except for one filament segment, indicate that the sign of the dominant helicity in the AR was negative. Global magnetic field extrapolations, before any AR activity, agree with this finding (Section~\\ref{S-Morph}). This is in strong contradiction with what is expected by models of MC formation constrained by the helicity conservation principle. In view of this result, we analyze the photospheric magnetic field evolution of the region and we perform an in-depth analysis of the local helicity flux injection along 18 November (Section~\\ref{S-LocHInj}). We find that a zone at the south of AR 10501 was the location of positive helicity flux injection during that day. Finally, we conclude that the ejected flux rope, observed later as an MC, should be located at this southern portion of the AR, and discuss the implications of our finding for CMEs/MCs triggering models (Section~\\ref{S-Concl}). ",
        "conclusions": "% \\label{S-Concl}  We have determined the global and local magnetic helicity sign of AR 10501 on 18 November 2003, using different methods based on the analysis of a multi-wavelength data set. We have also discussed the association of this AR with the MC of 20 November 2003, observed by ACE, the largest geoeffective cloud during the solar cycle 23.  AR 10501 is surrounded by a large apparent circular-shaped single filament, which is in fact formed by several distinct segments. The flares of  18 November 2003 were initiated by the continuous emergence of magnetic bipoles and by the eruption of some segments of the filament.  The destabilization of one of the filament segments is the primary trigger of the third flare. As discussed in the `CSHKP' (\\opencite{Carmichael64}; \\opencite{Sturrock66}; \\opencite{Hirayama74}, and \\opencite{Kopp76}; also Forbes and Malherbe (1991)) standard solar-flare model, a filament eruption (in our particular case, the eruption of a filament segment) due to some magnetic instability occurs above the magnetic inversion line. As a result, filament material moves away from the solar surface, the pre-existing magnetic arcade stretches upward, and the condition arises for magnetic reconnection, leading to a CME ejection and a subsequent MC in the interplanetary medium.  Based on the multi-wavelength data set, we have found the following results in relation to the magnetic helicity of the AR:  \\smallskip  \\noindent $-$ The very extended filament within and surrounding the AR presents evidence of dextral, as well as sinistral chirality, {\\it i.e.} negative and positive magnetic helicity.  \\noindent $-$ The sunspot whorls show a left-hand twist, which corresponds to negative magnetic helicity.  \\noindent $-$ The reverse `J-shaped' flare ribbons indicate negative helicity in the active region. However, this may be a false indicator because the ribbons follow the magnetic inversion line.  \\noindent $-$ From the coronal magnetic field model, we have found that the best $\\alpha$ values to fit the large-scale TRACE loops are negative. These negative $\\alpha$ values mean that the magnetic helicity is globally negative. However, at the south of the AR, where N4/P4 lie, the value of $\\alpha$ needed to locally match the shape of loops in the `post-flare' arcade after the flare that peaked at 08:31 UT has a tendency to be positive towards the east and negative towards the west, in agreement with the location of sinistral and dextral filament segments.  \\noindent $-$ The computation of the magnetic helicity injection, using G$_\\theta$ maps,  indicates that the main spot in the AR is dominated by negative helicity. However, there is a strong local positive injection of helicity in the southern polarities (N4/P4) of the AR.  \\smallskip From the above evidences, we conclude that AR 10501 has a global negative magnetic helicity. Despite this global negative helicity, the helicity density maps, {\\it i.e.} G$_\\theta$ maps, show a strong injection of positive magnetic helicity in the southern polarities. Our result provides a clear example of an AR in which the magnetic helicity sign is mixed, with simultaneous injection of both helicity signs. Previous works (\\opencite{Pevtsov97}; \\opencite{Green02}) have already reported examples in which the total helicity sign of an AR changed as it evolved because of parasitic magnetic polarity emergence having an opposite helicity sign to the main one. This is the first time that helicity flux density maps bring new information about the local helicity injection into an active region whose sign is opposite to the global helicity of the AR. Similar local injection of helicity of an opposite sign may explain the few cases of discrepancy found by \\inlinecite{Leamon04} between the helicity sign carried by MCs and the global helicity sign of their identified solar source region.  We speculate that due to a global instability in the AR, a flux rope with positive helicity erupted. This flux rope would contain the filament segment with positive helicity. As discussed above, the erupting filament has sections with opposite chiralities (sinistral in F2a and dextral in F2b). Filaments with the same chirality can merge, while those with different chirality cannot merge (\\opencite{Martin98}; \\opencite{Rust01}; \\opencite{Schmieder04}; \\opencite{Aulanier06}). The segment F2 corresponds to two different flux tubes having the same axial magnetic field direction in agreement with the counterstreaming motions observed before the eruption. However, the two sections F2a and F2b have opposite chiralities. Possibly, during the third flare, these two segments interacted through magnetic reconnection and helicity carried by both section partially canceled. The positive helicity carried by F2a being larger, the CME and ICME induced by the eruption transported away a net positive helicity, as measured by ACE at 1 AU \\cite{Gopalswamy05}.  Finally, our finding of a local injection of helicity with opposite signs is potentially important for eruption models. \\inlinecite{Kusano04} developed a model based on the reconnection of magnetic flux ropes with opposite helicity. Helicity annihilation potentially allows the release of a larger amount of free energy since the field can eventually relax to a state closer to potential.  \\begin{acks} The authors thanks Dr. Pascal D\\'emoulin for fruitful discussions and suggestions. R.C. thanks the Le Centre Franco-Indien pour la Promotion de la Recherche Avanc\\'ee (CEFIPRA) for his postdoctoral grant. This work was done in the frame of the European network SOLAIRE. This work used the DAVE/DAVE4VM codes written and developed by the Naval Research Laboratory. E.P. wishes to thank Peter Schuck for providing the DAVE/DAVE4VM code and for useful discussions. The work of E.P. was supported, in part, by the NASA HTP and SR$\\&$T programs. C.H.M. thanks the Argentinean grants: UBACyT X127 and PICT 03-33370 (ANPCyT). C.H.M. is a member of the Carrera del Investigador Cient\\'{i}fico, CONICET. We acknowledge the use of TRACE data. MDI data are a courtesy of SOHO/MDI consortium. SOHO is a project of international cooperation between ESA and NASA. B.S and C.H.M have started this work in the frame of the ISSI workshop chaired by Dr. Consuelo Cid (2008-2010). We also thank the anonymous referee for helpful and constructive comments. \\end{acks}"
    },
    "0910/0910.4529_arXiv.txt": {
        "abstract": "{Determination of the mass functions of open clusters of different ages allows us to infer the efficiency with which brown dwarfs are evaporated from clusters to populate the field.} {In this paper we present the results of a photometric survey to identify low mass and brown dwarf members of the old open cluster Praesepe (age 590$^{+150}_{-120}$\\,Myr, distance 190$^{+6.0}_{-5.8}$\\,pc) from which we estimate its mass function and compare this with that of other clusters.} {We performed an optical ($I_{\\rm c}$-band) and near-infrared ($J$ and $K_{\\rm s}$-band) photometric survey of Praesepe covering 3.1\\,deg$^2$. With 5$\\sigma$ detection limits of $I_{\\rm c}=23.4$ and $J=20.0$, our survey is predicted to be sensitive to objects with masses from 0.6 to 0.05\\,M$_\\odot$.} {We photometrically identify 123 cluster member candidates based on dust-free atmospheric models and 27 candidates based on dusty atmospheric models.  The mass function rises from 0.6\\,M$_\\odot$ down to 0.1\\,M$_\\odot$ (a power law fit of the mass function gives $\\alpha$=1.8$\\pm$0.1; \\,$\\xi$(M)\\,$\\propto$\\,M$^{-\\alpha}$\\,), and then turns over at $\\sim$0.1\\,M$_\\odot$. This rise agrees with the mass function inferred by previous studies, including a survey based on proper motion and photometry. In contrast, the mass function differs significantly from that measured for the Hyades, an open cluster with a similar age ($\\tau$\\,$\\sim$\\,600\\,Myr). Possible reasons are that the clusters did not have the same initial mass function, or that dynamical evolution (e.g.\\ evaporation of low mass members) has proceeded differently in the two clusters. Although different binary fractions could cause the observed (i.e.\\ system) mass functions to differ, there is no evidence for differing binary fractions from measurements published in the literature.  Of our cluster candidates, six have masses predicted to be equal to or below the stellar/substellar boundary at 0.072\\,M$_\\odot$.} {} ",
        "introduction": "Several publications in the past decade have been concerned with the mass function (MF) of low mass stellar and substellar populations in open clusters, including $\\sigma$~Orionis (\\citealt{bejar2002}, \\citealt{caballero2007}), the Orion Nebula Cluster (\\citealt{hillenbrand2000}, \\citealt{slesnick2004}), IC~2391 (\\citealt{barrado2004}, \\citealt{boudreault2009}), the Pleiades (\\citealt{moraux2003}, \\citealt{lodieu2007}), and the Hyades (\\citealt{reid99}, \\citealt{bouvier2008}), to name just a few. Studies of relatively old open clusters (age\\,$\\gtrsim$\\,100\\,Myr) are important for the following two reasons in particular. First, they permit a study of the intrinsic evolution of brown dwarfs (BDs), e.g.\\ their luminosity and effective temperature, which constrains structural and atmospheric models. Second, together with younger clusters we can investigate how BD populations as a whole evolve and thus probe the efficiency with which BDs evaporate from clusters to populate the Galactic field.  Numerical simulations of cluster evolution have demonstrated that the MFs can evolve through dynamical interaction (\\citealt{marcos2000}; \\citealt{adams2002b}).  These interactions result in a decrease of the open cluster BD (and low-mass star) population.  This has been observed by \\cite{bouvier2008} from a comparison of the Pleiades (120\\,Myr) and Hyades (625\\,Myr) mass functions.  Many earlier studies of the substellar MF have focused on young open clusters with ages less than $\\sim$\\,100\\,Myr, and in many cases much younger ($< 10$\\,Myr). This is partly because BDs are bright when they are young (lacking a significant nuclear energy source, they cool as they age), thus easing detection of the least massive objects. However, youth presents difficulties. First, intra-cluster extinction plagues the determination of the intrinsic luminosity function from the measured photometry. Second, at these ages the BD models have large(r) uncertainties (\\citealt{baraffe2002}). Estimates of the substellar MF in very young clusters (age $\\lesssim$1\\,Myr) might be unreliable due to these modelling uncertainties (\\citealt{chabrier2005}). BDs in older clusters suffer less from these problems, but have the disadvantage that much deeper surveys are required to detect them.  The old open cluster Praesepe is an interesting target considering its age and distance. It is located at a distance of 190$^{+6.0}_{-5.8}$\\,pc (based on parallax measurements from the new Hipparcos data reduction, \\citealt{Leeuwen2009}) and has an age of 590$^{+150}_{-120}$\\,Myr (by isochrone fitting in the Hertzsprung-Russell diagram; \\citealt{fossati2008}). The extinction towards this cluster is low, $E(B-V)$\\,=\\,0.027$\\pm$0.004\\,mag (\\citealt{taylor2006}), while determinations of the metallicity of Praesepe yield some discrepancies: [Fe/H]\\,=\\,0.038$\\pm$0.039, \\cite{friel1992}; +0.13$\\pm$0.10, \\cite{boesgaard1988}; 0.11$\\pm$0.03 from spectroscopy and 0.20$\\pm$0.04 from photometry, \\cite{an2007}; +0.27$\\pm$0.10, \\cite{pace2008}.  \\cite{hambly1995} presented a $\\sim$19\\,deg$^2$ survey of the Praesepe cluster down to masses of $\\sim$0.1\\,M$_\\odot$ and observed a rise of the MF at the lowest masses. They concluded that this implied a large population of BDs. A shallow survey complete to $I$\\,=\\,21.2\\,mag, $R$\\,=\\,22.2\\,mag over 800\\,arcmin$^2$ uncovered one spectrally confirmed very low-mass star or BD (spectral type of M8.5V) with a model-dependent mass of 0.063--0.084\\,M$_\\odot$ (\\citealt{Ma_98}). A survey over the central 1\\,deg$^2$ with 10$\\sigma$ limits of $R$\\,=\\,21.5, $I$\\,=\\,20.0 and $Z$\\,=\\,21.5\\,mag revealed 19~BD candidates and the first MF determination of Praesepe down to the substellar limit, but without spectral confirmation (\\citealt{pinfield97}). Subsequent infrared photometry of the sample reduced this number to nine candidates (\\citealt{hodgkin99}).  \\cite{adams2002a} presented a 100\\,deg$^2$ study of Praesepe using 2MASS (Two-Micron All Sky Survey) data and Palomar survey photographic plates, from which they derived proper motions. They determined the radial profile of this cluster but their MF does not reach the substellar regime.  A more recent proper motion survey of Praesepe covers a much larger area (300\\,deg$^2$; \\citealt{kraus2007}), but does not reach the BD regime either (the limit is $\\sim$0.12\\,M$_\\odot$). Finally, the most recent substellar MF determination of Praesepe was published by \\cite{gonzalez-garcia2006} and extends to a 5$\\sigma$ detection limit of $i=24.5$\\,mag corresponding to 0.050--0.055\\,M$_\\odot$.  They identified one new substellar candidate, but their survey covers only 1177 arcmin$^2$.  In this paper, we present the results of a program to study, in detail, the MF of Praesepe down to the substellar regime. Our photometric survey is, as with \\cite{gonzalez-garcia2006}, the deepest so far in optical and near-infrared (NIR) bands, with 5$\\sigma$ detection limits of $I_{\\rm c}=23.4$ and $J=20.0$ (corresponding to a mass limit of about 0.05\\,M$_\\odot$), but covers more than nine times the area.  Our paper is structured as follows. We first present the data set, reduction procedure and calibration in section \\ref{obs-data-calib}.  We then discuss our candidate selection procedure in section \\ref{selection} and the survey results in section \\ref{results-survey} before discussing the derived MF in section \\ref{mf-variation}. We conclude in section \\ref{conclusions}. ",
        "conclusions": "Conclusions}  We have presented the results of a survey to study the mass function of the old open cluster Praesepe. The survey consisted of optical $I_{\\rm c}$-band photometry and NIR $J$ and $K_{\\rm s}$-band photometry with a total coverage of 3.1\\,deg$^2$, down to the substellar regime, with a 5$\\sigma$ detection limit corresponding to 0.05\\,M$_\\odot$ (the detection completeness to this level is $\\sim$87\\%).  Our final sample yields 123 photometric cluster member candidates based on a selection assuming a dust-free atmosphere and 27 photometric cluster candidates based on a selection assuming a dusty atmosphere. We estimate the contamination by field M-dwarfs to be 13\\% or less. Among our cluster candidates, six objects have theoretical masses equal to or less than the stellar/substellar boundary at 0.072\\,M$_\\odot$.  We observed that the MF of Praesepe is characterized by a rise in the number of objects from 0.6\\,M$_\\odot$ down to 0.1\\,M$_\\odot$, followed by a turn-over in the MF at $\\sim$0.1\\,M$_\\odot$. The rise is in agreement with the Praesepe MFs derived in several previous studies (\\citealt{hambly1995}; \\citealt{kraus2007}; \\citealt{baker2009}) but disagrees with \\cite{adams2002a}.  We have compared the mass function of Praesepe with one derived for the Hyades and have observed a significant difference: while the Hyades has a maximum at 0.35\\,M$_\\odot$, Praesepe has a maximum at a much lower mass, 0.1\\,M$_\\odot$. Assuming that they have similar ages (as main sequence fitting suggests), we conclude that the clusters either had different {\\em initial} mass functions or that dynamical interaction has modified the evolution of one or both. More specifically, in the latter case, dynamical evaporation does not seem to have influenced the Hyades and Praesepe in the same way. A difference in the binary fraction or mass ratios could also cause a difference in the mass functions, but determinations of these are not yet precise enough to suggest any difference."
    },
    "0910/0910.4659_arXiv.txt": {
        "abstract": "The BEST wide-angle telescope installed at the Observatoire de Haute-Provence and operated in remote control from Berlin by the Institut f\\\"ur Planetenforschung, DLR, has observed the  CoRoT target fields prior to the mission. The resulting archive of stellar photometric lightcurves is used to search for deep transit events announced during CoRoT's alarm-mode to aid in fast photometric confirmation of these events. The \"initial run\" field of CoRoT (IRa01) has been observed with BEST in November and December 2006 for 12 nights. The first \"long run\" field (LRc01) was observed from June to September 2005 for 35 nights. After standard CCD data reduction, aperture photometry has been performed using the ISIS image subtraction method. About 30,000 lightcurves were obtained in each field. Transits of the first detected planets by the CoRoT mission, CoRoT-1b and CoRoT-2b, were found in archived data of the BEST survey and their lightcurves are presented here. Such detections provide useful information at the early stage of the organization of follow-up observations of satellite alarm-mode planet candidates. In addition, no period change was found over $\\sim$4 years between the first BEST observation and last available transit observations. ",
        "introduction": "The Berlin Exoplanet Search Telescope (BEST) is a 19.5 cm aperture wide angle telescope dedicated to time-series photometric observations \\cite{Rauer2004}. The main purpose of the instrument is to provide ground-based support to the CoRoT space mission (CNES; \\cite{Baglin2006}). The mission target fields are observed at least one year prior to the observations of the spacecraft. Therefore, planetary transit candidates found by CoRoT at bright stars  can be searched for quickly in the BEST archive to confirm the transit event and to check the ephemeris. In addition, BEST data sets are used to search for new variable stars in the CoRoT target fields which are input to the additional science program, like e.g. eclipsing binary stars (\\cite{Karoff2007}, \\cite{Kabath07}, \\cite{Kabath08}). Furthermore, information on variable stars nearby CoRoT transit candidates can be used to help disentangling crowding problems during the CoRoT lightcurve analysis.   The transit signals of the first two planets detected by CoRoT (CoRoT-1b and CoRoT-2b) (\\cite{barge2008}, \\cite{alonso2008}), were found using the \"alarm detection mode\" of the mission, during ongoing observations of the respective target fields. Since both planets orbit relatively bright stars, we searched for signatures of the transit candidates in our BEST data archive. Partial transit events of the planets were recorded by BEST in December 2006 and summer 2005 during observations of the CoRoT \"initial run\" and first \"long run\" fields. In the following, we describe these pre-discovery observations of CoRoT-1b and CoRoT-2b. ",
        "conclusions": "We present pre-discovery observations of the first two planets detected by the CoRoT space mission, CoRoT-1b and CoRoT-2b. The transit events were detected in the BEST data archive, based on the epochs determined from the CoRoT data (\\cite{barge2008}, \\cite{alonso2008}). Although the observational duty cycle was severely affected by bad weather conditions, partial transits of both planets were detected by BEST. A transit of CoRoT-1b was observed on 10 December 2006. Transits of CoRoT-2b were observed in 14 July 2005, 28 July 2005 and 04 Aug 2005. No significant O-C deviation in comparison to the ephemerides of \\cite{alonso2008} was found.  When planetary candidates are first announced from the CoRoT alarm-mode, their ephemeris can be based on few transit events only. At this point, confirmation from ground-based observations taken one or more years prior to the spacecraft observations aids in checking the ephemeris quickly. In addition, these observations confirm the transit on the prime target and help identifying close possibly contaminating variable stars. We will continue to use the BEST data archive for this purpose in future to support the planet search in particular during the alarm-mode of CoRoT, when primarily deep transits around bright stars are detected which overlap with the detection range of BEST."
    },
    "0910/0910.2887_arXiv.txt": {
        "abstract": "{The findings of more than 350 extrasolar planets, most of them nontransiting Hot Jupiters, have revealed correlations between the metallicity of the main-sequence (MS) host stars and planetary incidence. This connection can be used to calculate the planet formation probability around other stars, not yet known to have planetary companions. Numerous wide-field surveys have recently been initiated, aiming at the transit detection of extrasolar planets in front of their host stars. Depending on instrumental properties and the planetary distribution probability, the promising transit locations on the celestial plane will differ among these surveys.} {We want to locate the promising spots for transit surveys on the celestial plane and strive for absolute values of the expected number of transits in general. Our study will also clarify the impact of instrumental properties such as pixel size, field of view (FOV), and magnitude range on the detection probability.} {We used data of the Tycho catalog for $\\approx~1$ million objects to locate all the stars with $0^\\mathrm{m}~\\lesssim~m_\\mathrm{V}~\\lesssim~11.5^\\mathrm{m}$ on the celestial plane. We took several empirical relations between the parameters listed in the Tycho catalog, such as distance to Earth, $m_\\mathrm{V}$, and $(B-V)$, and those parameters needed to account for the probability of a star to host an observable, transiting exoplanet. The empirical relations between stellar metallicity and planet occurrence combined with geometrical considerations were used to yield transit probabilities for the MS stars in the Tycho catalog. Magnitude variations in the FOV were simulated to test whether this fluctuations would be detected by BEST, XO, SuperWASP and HATNet.} {We present a sky map of the expected number of Hot Jupiter transit events on the basis of the Tycho catalog. Conditioned by the accumulation of stars towards the galactic plane, the zone of the highest number of transits follows the same trace, interrupted by spots of very low and high expectation values. The comparison between the considered transit surveys yields significantly differing maps of the expected transit detections. While BEST provides an unpromising map, those for XO, SuperWASP, and HATNet show FsOV with up to 10 and more expected detections. The sky-integrated magnitude distribution predicts 20 Hot Jupiter transits with orbital periods between 1.5\\,d and 50\\,d and $m_\\mathrm{V}~<~8^\\mathrm{m}$, of which two are currently known. In total, we expect 3412 Hot Jupiter transits to occur in front of MS stars within the given magnitude range. The most promising observing site on Earth is at latitude $~=~-1$.} {} ",
        "introduction": "\\label{sec:intro}  A short essay by Otto Struve \\citep{1952Obs....72..199S} provided the first published proposal of transit events as a means of exoplanetary detection and exploration. Calculations for transit detection probabilities \\citep{1971Icar...14...71R, 1984Icar...58..121B, 2006AcA....56..183P} and for the expected properties of the discovered planets have been done subsequently by many others \\citep{2005A&A...442..731G, 2007A&A...475..729F, 2008ApJ...686.1302B}. Until the end of the 1990s, when the sample of known exoplanets had grown to more than two dozen \\citep{2000ApJ...532L..51C}, the family of so-called `Hot Jupiters', with 51 Pegasi as their prototype, was unknown and previous considerations had been based on systems similar to the solar system. Using geometrical considerations, \\citet{1971Icar...14...71R}\\footnote{A correction to his Eq. (2) is given in \\citet{1984Icar...58..121B}.} found that the main contribution to the transit probability of a solar system planet would come from the inner rocky planets. However, the transits of these relatively tiny objects remain undetectable around other stars as yet.  The first transit of an exoplanet was finally detected around the sun-like star HD209458 \\citep{2000ApJ...529L..45C, 2000A&A...359L..13Q}. Thanks to the increasing number of exoplanet search programs, such as the ground-based Optical Gravitational Lensing Experiment (OGLE) \\citep{1992AcA....42..253U}, the Hungarian Automated Telescope (HAT) \\citep{2002PASP..114..974B, 2004PASP..116..266B}, the Super Wide Angle Search for Planets (SuperWASP) \\citep{2003ASPC..294..405S}, the Berlin Exoplanet Search Telescope (BEST) \\citep{2004PASP..116...38R}, XO \\citep{2005PASP..117..783M}, the Transatlantic Exoplanet Survey (TrES) \\citep{2007ASPC..366...13A}, and the Tautenburg Exoplanet Search Telescope (TEST) \\citep{2009IAUS..253..340E} and the space-based missions `Convection, Rotation \\& Planetary Transits' (CoRoT) \\citep{2002ESASP.485...17B} and Kepler \\citep{2007CoAst.150..350C}, the number of exoplanet transits has grown to 62 until September $1^\\mathrm{st}$ 2009\\footnote{Extrasolar Planets Encyclopedia (EPE): www.exoplanet.eu. Four of these 62 announced transiting planets have no published position.} and will grow drastically within the next years. These transiting planets have very short periods, typically $<~10$\\,d, and very small semimajor axes of usually $<~0.1$\\,AU, which is a selection effect based on geometry and Kepler's third law \\citep{1619QB41.K38.......}. Transiting planets with longer periods present more of a challenge, since their occultations are less likely in terms of geometrical considerations and they occur less frequently.  Usually, authors of studies on the expected yield of transit surveys generate a fictive stellar distribution based on stellar population models. \\citet{2007A&A...475..729F} use a Monte-Carlo procedure to synthesize a fictive stellar field for OGLE based on star counts from \\citet{2006AcA....56....1G}, a stellar metallicity distribution from \\citet{2004A&A...418..989N}, and a synthetic structure and evolution model of \\citet{2003A&A...409..523R}. The metallicity correlation, however, turned out to underestimate the true stellar metallicity by about 0.1\\.dex, as found by \\citet{2004A&A...415.1153S} and \\citet{2005ApJ...622.1102F}. In their latest study, \\citet{2009A&A...504..605F} first generate a stellar population based on the Besan\\c{c}on catalog from \\citet{2003A&A...409..523R} and statistics for multiple systems from \\citet{1991A&A...248..485D} to apply then the metallicity distribution from \\citet{2004A&A...415.1153S} and issues of detectability \\citep{2006MNRAS.373..231P}. \\citet{2008ApJ...686.1302B} rely on a Galactic structure model by \\citet{1980ApJS...44...73B}, a mass function as suggested by \\citet{2002AJ....124.2721R} based on Hipparcos data, and a model for interstellar extinction to estimate the overall output of the current transit surveys TrES, XO, and Kepler. In their paper on the number of expected planetary transits to be detected by the upcoming Pan-STARRS survey \\citep{2004SPIE.5489...11K}, \\citet{2009A&A...494..707K} also used a Besan\\c{c}on model as presented in \\citet{2003A&A...409..523R} to derive a brightness distribution of stars in the target field and performed Monte-Carlo simulations to simulate the occurrence and detections of transits. These studies include detailed observational constraints such as observing schedule, weather conditions, and exposure time and issues of data reduction, e.g. red noise and the impact of the instrument's point spread function.  In our study, we rely on the extensive data reservoir of the Tycho catalog instead of assuming a stellar distribution or a Galactic model. We first estimate the number of expected exoplanet transit events as a projection on the complete celestial plane. We refer to recent results of transit surveys such as statistical, empirical relationships between stellar properties and planetary formation rates. We then use basic characteristics of current low-budget but high-efficiency transit programs (BEST, XO, SuperWASP, and HATNet), regardless of observational constraints mentioned above, and a simple model to test putative transits with the given instruments. With this procedure, we yield sky maps, which display the number of expected exoplanet transit detections for the given surveys, i.e. the transit sky as it is seen through the eyeglasses of the surveys.  The Tycho catalog comprises observations of roughly 1 million stars taken with the Hipparcos satellite between 1989 and 1993 \\citep{1997yCat.1239....0E, 1997ESASP.402...25H}. During the survey, roughly 100 observations were taken per object. From the derived astrometric and photometric parameters, we use the right ascension ($\\alpha$), declination ($\\delta$), the color index $(B-V)$, the apparent visible magnitude $m_\\mathrm{V}$, and the stellar distance $d$ that have been calculated from the measured parallax. The catalog is almost complete for the magnitude limit $m_\\mathrm{V}~\\lesssim~11.5^\\mathrm{m}$, but we also find some fainter stars in the list. ",
        "conclusions": "\\label{sec:discussion}  Our values for the XO project are much higher than those provided by \\citet{2008ApJ...686.1302B}, who also simulated the expected exoplanet transit detections of XO. This is due to their much more elaborate inclusion of observational constraints such as observational cadence, i.e. hours of observing per night, meteorologic conditions, exposure time, and their approach of making assumptions about stellar densities and the Galactic structure instead of using catalog-based data as we did. Given these differences between their approach and ours, the results are not one-to-one comparable. While the study of \\citet{2008ApJ...686.1302B} definitely yields more realistic values for the expected number of transit detections considering all possible given conditions, we provide estimates for the celestial distribution of these detections, neglecting observational aspects.  In addition to the crucial respects that make up the efficiency of the projects, as presented in Table \\ref{tab:surveys}, SuperWASP and HATNet benefit from the combination of two observation sites and several cameras, while XO also takes advantage of twin lenses but a single location. Each survey uses a single camera type and both types have similar properties, as far as our study is concerned. The transit detection maps in Fig. \\ref{fig:surveys} refer to a single camera of the respective survey. The alliance of multiple cameras and the diverse observing strategies among the surveys \\citep{2005PASP..117..783M, 2009IAUS..253...29C} bias the speed and efficiency of the mapping procedure. This contributes to the dominance of SuperWASP (18 detections, 14 of which have published positions)\\footnote{EPE as of September $1^\\mathrm{st}$ 2009} over HATNet (13 detections, all of which have published positions)\\footnotemark[3], XO (5 detections, all of which have published positions)\\footnotemark[3], and BEST (no detection)\\footnotemark[3].  It is inevitable that a significant fraction of unresolved binary stars within the Tycho data blurs our results. The impact of unresolved binaries without physical interaction, which merely happen to be aligned along the line of sight, is significant only in the case of extreme crowding. As shown by \\citet{2007ASPC..366..283G}, the fraction of planets not detected because of blends is typically lower than 10\\,\\%. The influence of unresolved physical binaries will be higher. Based on the empirical period distribution for binary stars from \\citet{1991A&A...248..485D}, \\citet{2008ApJ...686.1302B} estimate the fraction of transiting planets that would be detected despite the presence of binary systems to be $\\approx~70$\\,\\%. Both the contribution of binary stars aligned by chance and physically interacting binaries result in an overestimation of our computations of $\\approx~40$\\,\\%, which is of the same order as uncertainties arising from the empirical relationships we use. Moreover, as \\citet{2006MNRAS.367.1103W} have shown, the density of eclipsing stellar binary systems increases dramatically towards the Galactic center. To control the fraction of false alarms, efficient data reduction pipelines, and in particular data analysis algorithms, are necessary \\citep{2006MNRAS.365..165S}.  Recent evidence for the existence of ultra-short period planets around low-mass stars \\citep{2009IAUS..253...45S}, with orbital periods $~<~1$\\,d, suggests that we underestimated the number of expected transits to occur, as presented in Sect. \\ref{sub:occurance}. The possible underestimation of exoplanets occurring at $\\mathrm{[Fe/H]}_\\star~<~0$ also contributes to a higher number of transits and detections than we computed here. Together with the fact that the Tycho catalog is only complete to $m_\\mathrm{V} \\lesssim 11.5^\\mathrm{m}$, whereas the surveys considered here are sensitive to slightly fainter stars (see Table \\ref{tab:surveys}), these trends towards higher numbers of expected transit detections might outweigh the opposite effect of unresolved binary stars.  A radical refinement of both our maps for transits occurrence and detections will be available within the next few years, once the `Panoramic Survey Telescope and Rapid Response System' (Pan-STARRS) \\citep{2002SPIE.4836..154K} will run to its full extent. Imaging roughly 6000 square degrees every night with a sensitivity down to $m_\\mathrm{V} \\approx 24$, this survey will not only drastically increase the number of cataloged stars -- thus enhance our knowledge of the localization of putative exoplanetary transits -- but could potentially detect transits itself \\citep{2009ApJ...704.1519D}. The Pan-STARRS catalog will provide the ideal sky map, on top of which an analysis presented in this paper can be repeated for any ground-based survey with the aim of localizing the most appropriate transit spots on the celestial plane. The bottleneck for the verification of transiting planets, however, is not the localization of the most promising spots but the selection of follow-up targets accessible with spectroscopic instruments. The advance to fainter and fainter objects thus won't necessarily lead to more transit confirmations. Upcoming spectrographs, such as the ESPRESSO{\\MVAt}VLT and the CODEX{\\MVAt}E-ELT \\citep{2008PhST..130a4007P}, can be used to confirm transits around fainter objects. These next-generation spectrographs that will reveal Doppler fluctuations on the order of cm$\\cdot$s$^{-1}$ will also enhance our knowledge about Hot Neptunes and Super-Earths, which the recently discovered transits of GJ\\,436\\,b \\citep{2004ApJ...617..580B}, HAT-P-11\\,b \\citep{2009arXiv0901.0282B}, and CoRoT-7b \\citep{2009arXiv0908.0241L} and results from \\citet{2009IAUS..253..502L} predict to be numerous.  Further improvement of our strategy will emerge from the findings of more exoplanets around MS stars and from the usage of public data reservoirs like the NASA Star and Exoplanet Database\\footnote{http://nsted.ipac.caltech.edu}, making assumptions about the metallicity distribution of planet host stars and the orbital period distribution of exoplanets more robust."
    },
    "0910/0910.0882_arXiv.txt": {
        "abstract": "\\noindent Globular  clusters are an  important test bed  for Newtonian gravity  in  the  weak-acceleration  regime,  which is  vital  to  our understanding of the nature  of the gravitational interaction.  Recent claims  have  been  made  that  the velocity  dispersion  profiles  of globular  clusters flatten  out at  large radii,  despite  an apparent paucity  of dark matter  in such  objects, indicating  the need  for a modification of gravitational theories.  We continue our investigation of this claim,  with the largest spectral samples  ever obtained of 47 Tucanae  and  M55.   Furthermore,  this  large sample  allows  for  an accurate metallicity calibration based on the equivalent widths of the calcium triplet  lines and $K$  band magnitude of  the Tip of  the Red Giant Branch.   Assuming an isothermal distribution,  the rotations of each  cluster are also  measured with  both clusters  exhibiting clear rotation signatures.   The global velocity dispersions of  NGC 121 and Kron 3, two globular clusters  in the Small Magellanic Cloud, are also calculated.   By applying  a simple  dynamical model  to  the velocity dispersion   profiles  of  47   Tuc  and   M55,  we   calculate  their mass-to-light profiles, total masses and central velocity dispersions. We  find  no  statistically  significant flattening  of  the  velocity dispersion at large radii for M55,  and a marked {\\it increase} in the profile of 47 Tuc for  radii greater than approximately half the tidal radius.   We  interpret this  increase  as  an evaporation  signature, indicating that 47 Tuc is undergoing, or has undergone, core-collapse, but  find  no  requirement  for  dark  matter  or  a  modification  of gravitational theories in either cluster. ",
        "introduction": "The  nature  of the  gravitational  interaction  is  one of  the  most important concepts in astrophysics, yet complete comprehension of this interaction  is  still  elusive.   The so-called  Pioneer  and  Fly-by anomalies, where spacecraft exhibit  behaviour that is unexpected from Newtonian  and general relativity  gravitation theories,  outline this lack       of       understanding      \\cite[see][and       references therein]{Anderson02,deDiego08}, although these  examples may have more mundane  explanations.  More  importantly,  it has  been claimed  that several globular  clusters (GCs; $\\omega$~Centauri, M15,  M30, M92 and M107) exhibit  a flattening of  their velocity dispersion  profiles at radii  $R\\sim\\frac{r_t}{2}$, where $r_t$  is the  tidal radius  of the cluster \\citep{Scarpa03,Scarpa04a,Scarpa04b}.   The authors claim that either dark matter (DM), or a modification of gravitational theory, is required to explain their results.  Modified theories of gravity  \\cite[MOG; see][for a review of modified gravity   theories]{Durrer08}   and   those  of   Newtonian   dynamics \\citep[MOND;][]{Milgrom83}  have  been  shown  to solve  some  of  the discrepancies. However,  these are not universal  theories and to-date have only  been applied to  specific instances \\cite[e.g.   the Bullet Cluster          and          galaxy         rotation          curves; see][respectively]{Angus06,Sanders07}.   MOG   theories  diverge  from Newtonian gravity  in the  high-acceleration regime and  MOND diverges from Newton  in the low-acceleration  regime. Therefore, if  either of these theories  were correct, the  effect should be measurable  at the predicted accelerations.  Independent of MOG or MOND theories, testing the gravitational interaction at low accelerations is essential to the overall understanding of gravity.  Globular  clusters   are  an  ideal  testing   ground  for  weak-field gravitation because  the accelerations  experienced by stars  at large radii are below the limit  where DM, or a modified gravitation theory, is  required  to  explain   observations  in  many  dynamical  systems \\cite[$a_0\\approx1.2\\times10^{-10}$\\,m\\,s$^{-2}$;][]{Scarpa07}. Furthermore, they are thought to contain little, or no, dark matter -- indicated   by  dynamical   models  \\citep{Phinney93},   {\\it  N}-body simulations   \\citep{Moore96},   observations   of  GC   tidal   tails \\citep{Odenkirchen01},   dynamical   and   luminous  masses   of   GCs \\citep{Mandushev91} and  the lack of microlensing  events from GC-mass dark  haloes \\citep{Navarro97,Ibata02}, although  this is  still under debate.  GCs are also located  at varying distances from the centre of the  Galaxy,  so  that  if  all exhibit  similar  behaviour,  Galactic influences cannot be the primary cause.  In \\citealt{Lane09} (hereafter  \\citetalias{Lane09}) we calculated the velocity dispersions  and mass-to-light profiles of M22,  M30, M53 and M68.   Our   conclusions  were  that  there  is   no  requirement  for significant quantities of dark  matter, or a modification of Newtonian gravity, to explain  the kinematics of any of  these clusters.  In the current  paper  we  continue   this  investigation  with  the  largest spectroscopic  dataset to date  of 47  Tucanae and  M55.  We  begin by describing the data acquisition/reduction (Section \\ref{data}) and the membership selection for each cluster (Section \\ref{membership}).  Our large samples allow us to determine a metallicity calibration based on the Tip  of the Red  Giant Branch (TRGB; Section  \\ref{metaldist}), as well  as the rotations  (Section \\ref{rotation}),  systemic velocities and    velocity   dispersions    (Section    \\ref{dispersions}),   and mass-to-light  profiles  (Section \\ref{ML})  --  where  we also  place limits  on  the  DM  content  of  each  cluster  from  their  velocity dispersions  and  mass-to-light  profiles.   Finally,  our  concluding remarks are presented in Section \\ref{conclusions}. ",
        "conclusions": "Having the largest sample of spectra ever obtained for 47 Tuc and M55, we were able to produce a very accurate calibration of [Fe/H] based on the equivalent  width of  the calcium triplet  lines and the  $K$ band magnitude  of   the  TRGB.   This   method  is  similar  to   that  by \\cite{Cole04} and \\cite{Warren09}, except that we use the TRGB instead of the  HB which  means this  method can  be used  for much  more distant objects.  We  calculated  the rotation  of  our  clusters  assuming them  to  be isothermal.   The  rotation  of  47  Tuc is  $\\sim2.2\\pm0.2$  with  an approximate   projected  rotational   axis   along  the   line  PA   = $40^\\circ-220^\\circ$,  and  M55  exhibits   rotation  at  a  level  of $0.25\\pm0.09$\\,km\\,s$^{-1}$  and has an  approximate axis  of rotation along the  line PA =  $65^\\circ-245^\\circ$.  For 47 Tuc,  the rotation amplitude    is    in    good    agreement    with    previous    work \\cite[e.g.][]{Meylan86}.   The  only  previous  study  estimating  the rotation of  M55 \\citep{Szekely07} found  a value about twice  that of the current work, but \\cite{Szekely07} had a sample size approximately half that of ours, which may explain this discrepancy.  Our calculated velocity dispersion profiles  of 47 Tuc and M55 provide no evidence that either DM  or a modification of current gravitational theories  are  required   to  reconcile  their  kinematic  properties, corroborating previous  work in \\citetalias{Lane09}.   The dynamics of M55 are  well described by a purely  analytic \\citet{Plummer11} model, which  indicates  that  Newtonian  gravity  adequately  describes  its velocity dispersions,  and shows no breakdown of  Newtonian gravity at $a_0\\approx1.2\\times10^{-10}$\\,ms$^{-2}$,  as  has  been  claimed  for other GCs.  The  internal dynamics of 47 Tuc  (for $R<r_t/2$) are also very  well  described by  the  Plummer  model,  however, the  velocity dispersion profile of 47 Tuc  exhibits a large increase for $R>r_t/2$, exactly the  region where evaporation due to  two-body interactions in the  core   should  be  observable,  especially   for  GCs  undergoing core-collapse  or in  a post-core-collapse  state.  We  interpret this increase in  velocity dispersion as  evaporation, and hence  that this cluster  is either  presently in  a state  of core-collapse,  or  is a post-core-collapse GC. This  adds to the growing evidence  that 47 Tuc is  currently  undergoing  a  dynamical phase  change  \\cite[e.g.][and references therein]{Gebhardt95b,Robinson95,Howell00}.  A full analysis of the  outer regions of  this apparently evaporating cluster  will be performed in a subsequent paper.  We  used a Plummer  model to  determine the  total mass,  scale radius ($r_s$), and the M/L$_{\\rm V}$ profile for each cluster.  We find that neither  cluster has  M/L$_{\\rm  V}\\gg1$, indicating  that  DM is  not dominant.   Within  the uncertainties,  our  estimated cluster  masses match those  in the  literature well, as  do the M/L$_{\\rm  V}$ ratios \\cite[e.g.][]{Meylan89,Pryor93,Meziane96,Kruijssen09}.              The mass-to-light  profiles  produced   by  \\cite{Gebhardt95a}  cannot  be compared to the current work  because they sampled the inner $10'$ for which the mass is uncertain. Note that we consider using mass-to-light {\\it profiles} is a more accurate method for calculating M/L$_{\\rm V}$ than  using the  core mass  and  surface brightness,  because of  this uncertainty.  While our results strongly  indicate that the current understanding of globular clusters being dark matter poor, and their dynamics explained by  standard Newtonian  gravity,  more robust  dynamical modelling  is required  for   confirmation."
    },
    "0910/0910.5469_arXiv.txt": {
        "abstract": "We analyze the statistics and star formation rate obtained in high-resolution numerical experiments of forced supersonic turbulence, and compare with observations. We concentrate on a systematic comparison of solenoidal (divergence-free) and compressive (curl-free) forcing (Federrath et~al.~2009~a,b), which are two limiting cases of turbulence driving. Our results show that for the same RMS Mach number, compressive forcing produces a three times larger standard deviation of the density probability distribution. When self-gravity is included in the models, the star formation rate is more than one order of magnitude higher for compressive forcing than for solenoidal forcing. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.0285_arXiv.txt": {
        "abstract": "% In this paper we discuss recent applications of the Smoothed Particle Hydrodynamics (SPH) method to the simulation of supersonic turbulence in the interstellar medium, as well as giving an update on recent algorithmic developments in solving the equations of magnetohydrodynamics (MHD) in SPH. Using high resolution calculations (up to 134 million particles), we find excellent agreement with grid-based results on a range of measures including the power spectrum slope in both the velocity field and the density-weighted velocity $\\rho^{1/3} v$, the latter showing a Kolmogorov-like $k^{-5/3}$ scaling as proposed by Kritsuk et al. (2007). We also find good agreement on the statistics of the Probability Distribution Function (PDF) and structure functions, independently confirming the scaling found by \\citet*{sfk08}. On Smoothed Particle Magnetohydrodynamics (SPMHD) we have recently wasted a great deal of time and effort investigating the vector potential as an alternative to the Euler potentials formulation, in the end concluding that using the vector potential has even more severe problems than the standard (${\\bf B}$-field based) SPMHD approach. ",
        "introduction": "Turbulence and magnetic fields are thought to be two of the most important ingredients in the star formation process, so much so that there remains ongoing debate --- both observational and theoretical --- as to which one is \\emph{the} controlling factor \\citep[e.g.][]{mk04,km09}. We have only recently begun to understand the role of either in depth, primarily as a result of our increased ability to simulate both in numerical calculations of the star formation process.   Several authors have proposed that turbulence in the interstellar medium (ISM) can be used to predict the form of the stellar initial mass function (IMF)  thus providing a `theory' of star formation \\citep{pn02,HennebelleChabrier2008}. Any such theory requires consensus on the basic statistical characteristics of turbulence in the supersonic regime and assumes that these are universal --- that is, independent of boundary conditions and driving mechanisms --- for which there is some observational support \\citep[e.g.][]{hb04}. However there exists considerable disagreement between results obtained using different numerical codes, most recently between \\citet{padoanetal07} who claim that the statistics of turbulence are universal and \\citet{bpetal06} who claim that they are not, based on calculations utilising both a Smoothed Particle Hydrodynamics (SPH) code and a grid-based Total-Variation-Diminishing (TVD) code.  This kind of disagreement, and the need within star formation theory for a consensus on the basic statistics of supersonic turbulence has prompted at least two major code comparison projects over the last few years, the ``Potsdam'' comparison \\citep{kitsionasetal09}, comparing decaying, hydrodynamic turbulence and the KITP comparison (unpublished), comparing decaying hydrodynamic and magnetohydrodynamic (MHD) turbulence. However these comparisons suffer from the limited time evolution that can be obtained from a decaying turbulence simulation as well as the numerical issues posed by starting from an evolved snapshot produced by a particular code.  We have therefore undertaken our own very detailed comparison of \\emph{driven} turbulence using just two codes, an SPH code, \\textsc{phantom}, and a grid code, \\textsc{flash} (used in uniform grid mode) taken to be broadly representative of their class of codes, in order to see whether agreement can be reached (and at what resolution) on the statistics of supersonic turbulence appropriate to the ISM. The results are published in full in \\citet{pf10} and we only provide a snapshot here of the main results (\\S\\ref{sec:turbulence}), referring the reader to that paper for more detailed information.  What we are not yet able to address is the combination of \\emph{driven} turbulence and MHD in SPH, primarily because of the limitations posed by the Euler potentials formulation used as the only method that sufficiently maintains the divergence-free ($\\nabla\\cdot{\\bf B} = 0$) constraint on the magnetic field such that star formation calculations can be performed \\citep[e.g.][]{pb07,pb08,pb09} without the stars themselves exploding due to numerical errors (see \\S\\ref{sec:mhd}). Following a suggestion from Axel Brandenburg, we therefore embarked on a quest to examine whether or not the use of the vector potential could similarly solve the divergence problem in the context of Smoothed Particle Magnetohydrodynamics (SPMHD) without the associated topological restrictions of the Euler potentials \\citep[see][]{brandenburg10}. In short, this turned out not to be the case \\citep[read the horror that is][]{price10}. However, for the pleasure of the reader we describe the tortuous journey in \\S\\ref{sec:mhd} ",
        "conclusions": ""
    },
    "0910/0910.2249_arXiv.txt": {
        "abstract": "We report the detection of $\\gamma$-ray pulsations ($\\ge 0.1$ GeV) from PSR~J2229+6114 and PSR~J1048$-$5832, the latter having been detected as a low-significance pulsar by EGRET. Data in the $\\gamma$-ray band were acquired by the Large Area Telescope aboard the \\textit{Fermi Gamma-ray Space Telescope}, while the radio rotational ephemerides used to fold the $\\gamma$-ray light curves were obtained using the Green Bank Telescope, the Lovell telescope at Jodrell Bank, and the Parkes telescope. The two young radio pulsars, located within the error circles of the previously unidentified EGRET sources 3EG~J1048$-$5840 and 3EG~J2227+6122, present spin-down characteristics similar to the Vela pulsar. PSR~J1048$-$5832 shows two sharp peaks at phases $0.15 \\pm 0.01$ and $0.57 \\pm 0.01$ relative to the radio pulse confirming the EGRET light curve, while PSR~J2229+6114 presents a very broad peak at phase $0.49 \\pm 0.01$. The $\\gamma$-ray spectra above 0.1 GeV of both pulsars are fit with power laws having exponential cutoffs near 3 GeV, leading to integral photon fluxes of $(2.19 \\pm 0.22 \\pm 0.32) \\times 10^{-7}$\\,cm$^{-2}$\\,s$^{-1}$  for PSR~J1048$-$5832 and $(3.77 \\pm 0.22 \\pm 0.44) \\times 10^{-7}$\\,cm$^{-2}$\\,s$^{-1}$ for PSR~J2229+6114. The first uncertainty is statistical and the second is systematic. PSR~J1048$-$5832 is one of two LAT sources which were entangled together as 3EG~J1048$-$5840. These detections add to the growing number of young $\\gamma$-ray pulsars that make up the dominant population of GeV $\\gamma$-ray sources in the Galactic plane. ",
        "introduction": "The nature of unidentified high-energy $\\gamma$-ray sources in the Galaxy was one of the major unanswered questions at the end of the EGRET era. The third EGRET catalog contained 170 unidentified sources, 74 of which were at Galactic latitude $|b|<10$\\degr \\citep{hartman99,bhatta03}. Rotation-powered pulsars are believed to dominate the Galactic $\\gamma$-ray source population (e.g. \\citet{yadiga95}), but their visibility is linked to their beam patterns. Soon after launch, the {\\it Fermi Gamma-ray Space Telescope} began to unveil many 3EG sources, discovering the radio-quiet pulsar in the CTA1 supernova remnant associated with 3EG~J0010+7309 \\citep{abdo08}, detecting the radio-loud pulsar PSR~J2021+3651 (associated with 3EG~J2021+3716) \\citep{abdo09b}, seen independently by {\\em AGILE} \\citep{halpern08}, as well as the radio pulsar PSR~J1028$-$5819, associated with 3EG~J1027$-$5817 \\citep{abdo09c} and new populations of radio-quiet $\\gamma$-ray pulsars, detectable using blind search techniques \\citep{abdo09f}. For the pulsars detected in $\\gamma$ rays, the bulk of the electromagnetic power output is in high energies. The $\\gamma$-ray emission is thus crucial for understanding the emission mechanism which converts the rotational energy of the neutron star into electromagnetic radiation. The discovery of many new $\\gamma$-ray pulsars will provide strong constraints on the location of the $\\gamma$-ray emitting regions, whether above the polar caps \\citep{daugherty96}, or far from the neutron star in the so-called ``outer gaps\" \\citep{romani96}, or in the intermediate regions like the ``slot gap\" \\citep{muslimov04} having ``two-pole caustic\" geometry \\citep{dyks03}. In this paper we report the \\textit{Fermi} detection of the two pulsars PSR~J1048$-$5832 and PSR~J2229+6114, which have spin characteristics similar to other young pulsars typified by the Vela pulsar. \\citet{kramer03} provide a discussion of ``Vela-like'' pulsars. \\\\  PSR~J1048$-$5832 (B1046$-$58) is located in the Carina region at low Galactic latitude ($l=287.42$\\degr, $b=0.58$\\degr). It was discovered during a 1.4 GHz Parkes survey of the Galactic plane and has a period  $P~\\sim$123.7~ms \\citep{johnston92}. It has a spin-down luminosity $\\dot{E}=4 \\pi^{2} I (\\dot{P}/P^{3})$ of $2\\times 10^{36}$\\,erg\\,s$^{-1}$, for a moment of inertia $I$ of $10^{45}$\\,g\\,cm$^{2}$, a surface dipole magnetic field strength of $3.5 \\times 10^{12}$\\,G, and a characteristic age $\\tau_{c} = P/2\\dot{P}$ of 20\\,kyr. High-resolution observations of the coincident $ASCA$ source by the \\textit{Chandra X-ray Observatory} and {\\it XMM-Newton\\/} revealed an asymmetric pulsar wind nebula (PWN) of $\\sim6''\\times11''$, surrounding a point source coincident with the pulsar but so far no X-ray pulsations were detected \\citep{gonzalez06}. PSR~J1048$-$5832 has been proposed as a counterpart of the steady EGRET source 3EG~J1048$-$5840 \\citep{fierro95,pivovaroff00,nolan03}, as suggested by positional coincidence, spectral, and energetic properties. Detailed analysis of the EGRET data found possible $\\gamma$-ray pulsations at $E>$400 MeV, a double-peaked light curve with $\\sim 0.4$ peak phase separation (Fig.\\ref{fig:J1048-5832_LC_multi_energy}, \\citealt{kaspi00}). An HI distance determination for the pulsar yields between 2.5 and 6.6 kpc \\citep{johnston96}. The NE2001 model \\citep{cordes02} assigns a distance of 2.7~kpc based partly on the HI distance determination. For this paper we adopt 3~kpc as the distance to the pulsar.\\\\  PSR J2229+6114 is located at ($l$,$b$) = (106.6\\degr,2.9\\degr) within the error box of the EGRET source 3EG~J2227+6122 \\citep{hartman99}. Detected as a compact X-ray source by \\textit{ROSAT} and \\textit{ASCA} observations of the EGRET error box, it was later discovered to be a radio and X-ray pulsar with a period of $P$ = 51.6 ms \\citep{halpern01b}. The radio pulse profile shows a single sharp peak, while the X-ray light curve at 0.8\\ --\\ 10 keV consists of two peaks, separated by $\\Delta\\phi = 0.5$. {\\em AGILE} recently reported the discovery of $\\gamma$-ray pulsations above 100 MeV \\citep{pellizzoni09}. The pulsar is as young as the Vela pulsar (characteristic age $\\tau_{c} = 10$\\,kyr), as energetic ($\\dot{E} = 2.2 \\times 10^{37}$\\,erg\\,s$^{-1}$), and is evidently the energy source of the ``Boomerang'' arc-shaped PWN G106.65+2.96, suggested to be part of the supernova remnant (SNR) G106.3+2.7 discovered by \\citet{joncas90}. Recently, the PWN has been detected at TeV energies by MILAGRO \\citep{abdo09g}. Studies of the radial velocities of both neutral hydrogen and molecular material place the system at $\\sim$ 800 pc \\citep{kothes01}, while \\citet{halpern01a} suggest a distance of 3 kpc estimated from its X-ray absorption. The pulsar DM, used in conjunction with the NE2001 model, yields a distance of 7.5 kpc\\,, significantly above all other estimates. For this paper we again adopt a distance of 3\\,kpc. ",
        "conclusions": "Although PSR~J1048$-$5832 and PSR~J2229+6114 are both Vela-like in age and spin characteristics, their light curves and derived emission geometries are quite different. Table \\ref{tab:sumres} summarises the main quantities measured for both pulsars. The double-peaked light curve of PSR J1048$-$5832 is nearly identical to that of Vela, whereas its derived efficiency is a factor of 10 larger than that of Vela \\citep{abdo09a}, adopting $f_{\\Omega} = 1$ and $d = 3$\\,kpc. On the contrary, the $\\gamma$-ray efficiency of J2229+6114 is very similar to that of the Vela pulsar, but the pulsar shows a single, large peak similar to PSR~J0357+32 discovered by searching for pulsations at the positions of bright $\\gamma$-ray sources \\citep{abdo09d}. Note that the efficiency of PSR~J2229+6114 would be smaller at the distance of about 1\\,kpc that some authors suggest.  \\begin{table*}[t] \\caption{This table summarises the results of the timing and spectral analysis of PSR~J1048$-$5832 and PSR J2229+6114. Statistical and systematics errors are reported.} \\vspace{0.1cm} \\begin{center}  \\begin{tabular}{llcc} \\hline Analysis & Parameters & PSR~J1048$-$5832 & PSR~J2229+6449  \\\\ \\hline \\hline \\textbf{Timing results} & Number of pulsed $\\gamma$ rays & 933 $\\pm$ 93 & 1365 $\\pm$ 97\\\\ & Peak position ($\\phi$)& 0.15 $\\pm$ 0.01 $\\pm$ 0.0001 (P1) & 0.49 $\\pm$ 0.01 $\\pm$ 0.001 \\\\ &              & 0.57 $\\pm$ 0.01 $\\pm$ 0.0001 (P2) & \\\\ & Peak separation ($\\Delta$) & 0.42 $\\pm$ 0.01 $\\pm$ 0.0001 & \\\\ & Peak FWHM& 0.06 $\\pm$ 0.01 (P1) & 0.23 $\\pm$ 0.03 \\\\ &          & 0.10 $\\pm$ 0.02 (P2) &  \\\\ \\hline \\textbf{Spectral results} & $^{a}$F (10$^{-7}$ cm$^{-2}$s$^{-1}$) & 2.19 $\\pm$ 0.22 $\\pm$ 0.32 & 3.77 $\\pm$ 0.22 $\\pm$ 0.44 \\\\ & $^{b}$F$_{E}$ (10$^{-11}$ erg cm$^{-2}$s$^{-1}$) & 19.4 $\\pm$ 1.0 $\\pm$ 3.1 & 23.7 $\\pm$ 0.7 $\\pm$ 2.5 \\\\ & $\\Gamma$& 1.38 $\\pm$ 0.06 $\\pm$ 0.12 & 1.85 $\\pm$ 0.06 $\\pm$ 0.10 \\\\ & $^{c}$E$_{c}$ (GeV)& 2.3$^{+0.3}_{-0.4}$ $\\pm$ 0.3 & 3.6$^{+0.9}_{-0.6}$ $\\pm$ 0.6 \\\\ & L$_{\\gamma}$ (10$^{35}$ erg s$^{-1}$) & 2.1 $f_{\\Omega}$ (d/3kpc)$^{2}$ & 2.6 $f_{\\Omega}$ (d/3kpc)$^{2}$ \\\\ & $\\eta_{\\gamma}$ & 0.10 $f_{\\Omega}$ (d/3kpc)$^{2}$& 0.011 $f_{\\Omega}$ (d/3kpc)$^{2}$ \\\\ \\hline \\end{tabular} \\tablenotetext{a}{Integral photon flux ($E>$0.1 GeV)} \\tablenotetext{b}{Integral energy flux ($E>$0.1 GeV)} \\tablenotetext{c}{Energy of an exponential cut-off to a power-law spectrum with index $\\Gamma$.} \\end{center} \\label{tab:sumres} \\end{table*}  With the growing number of detected $\\gamma$-ray pulsars, we are beginning to sample a wider variety of emission and viewing geometries, and pulsar ages. The range of light curve morphologies should allow improved constraints on high-energy emission models and a better understanding of the pulsar magnetospheric structure and acceleration process. For example, while many young pulsars, like J1048$-$5832, show Vela-type light curves, a small number are similar to J2229+6114, with a strong P2 component and a weak or absent P1 \\citep{welte09b}. Mapping the angle range over which P1 is missing, especially when viewing angle constraints are available, can help narrow down the high altitude emission zone. A larger pulsar sample also allows a study of evolution of the $\\gamma$-ray beaming and efficiency with pulsar age:~the pulsars seen with \\textit{Fermi}, not including millisecond pulsars, span ages from 10$^{3}$ to $2 \\times 10^{6}$\\,yr \\citep{abdo09i}, with hopes to extend that to larger $\\tau$ (lower $\\dot{E}$) as observations continue. Analysis of the population of pulsars with interpulses \\citep{welte08b} and radio polarization data \\citep{tauris98} have given hints that the magnetic inclination is larger for young pulsars and decreases with age. Thus we may even probe evolution of magnetic alignment during pulsar spindown."
    },
    "0910/0910.5005_arXiv.txt": {
        "abstract": "We measure shifts of the acoustic scale due to nonlinear growth and redshift distortions to a high precision using a very large volume of high-force-resolution simulations. We compare results from various sets of simulations that differ in their force, volume, and mass resolution. We find a consistency within $1.5-\\sigma$ for shift values from different simulations and derive shift $\\al(z) -1  = (0.300\\pm 0.015)\\% [D(z)/D(0)]^{2}$ using our fiducial set. We find a strong correlation  with a non-unity slope between shifts in real space and in redshift space and a weak correlation between the initial redshift and low redshift. Density-field reconstruction not only removes the mean shifts and reduces errors on the mean, but also tightens the correlations. After reconstruction, we recover a slope of near unity for the correlation between the real and redshift space and restore a strong correlation between the initial and the low redshifts. We derive propagators and mode-coupling terms from our \\Nb\\ simulations and compare with the \\Zel\\ approximation and the shifts measured from the $\\chi^2$ fitting, respectively. We interpret the propagator and the mode-coupling term of a nonlinear density field in the context of an average and a dispersion of its complex Fourier coefficients relative to those of the linear density field; from these two terms, we derive a signal-to-noise ratio of the acoustic peak measurement. We attempt to improve our reconstruction method by implementing 2LPT and iterative operations, but we obtain little improvement. The Fisher matrix estimates of uncertainty in the acoustic scale is tested using $5000\\trihGpc$ of cosmological PM simulations from \\citet{Taka09}. At an expected sample variance level of 1\\%, the agreement between the Fisher matrix estimates based on \\citet{SE07} and the \\Nb\\ results is better than 10 \\%. ",
        "introduction": "In recent years, attention to baryon acoustic oscillations (BAO) as a dark energy probe has increased unprecedentedly due to its robust nature against systematics, and they are now an essential component of most of the major future dark energy surveys under consideration. BAO originate from the sound waves that propagated through the hot plasma of photons and baryons in the very early Universe. At the epoch of recombination, photons and baryons decouple, and as a result, the sound waves freeze out, leaving a distinctive oscillatory feature in the large-scale structure of the cosmic microwave background \\citep[e.g.,][]{Mil99,deB00,Han00,Lee01,HalDasi,Netter02,Pearson03,BenoitArcheops,BennettWmap,Hinshaw07,Hinshaw08} and the matter density fields in Fourier space \\citep[e.g.,][]{Peebles70,SZ70,Bond84,Holtzman89,HS96,Hu96,EH98,Meiksin99}; In configuration space, the BAO appears as a single spherical peak at its characteristic scale.  The characteristic physical scale of this oscillatory feature, BAO, is the distance that the sound waves have traveled before the epoch of recombination, which is known as ``sound horizon scale''. This sound horizon scale is and will be measured precisely from current and future CMB data. With the knowledge of the physical scale, BAO can be used as a standard ruler to measure angular diameter distance and Hubble parameter at various redshifts and therefore provide critical information to identify dark energy \\citep[e.g.,][]{Hu96,Eisen03,Blake03,Linder03,Hu03,SE03}. Recently, BAO have been detected from large-scale structure of galaxy distributions and have been used to place an important constraint on dark energy \\citep[][]{Eisen05,Cole05,Hutsi06,Tegmark06,Percival07a,Percival07b,Blake07,Pad07,Okumura08,Estra08,Gazt08a,Gazt08b,Sanchez09,Percival09,Kazin09}.  Due to nonlinear structure growth at late times, the oscillatory feature of the BAO is increasingly damped, proceeding from small scales to larger scales, with decreasing redshift \\citep[e.g.,][]{Meiksin99,SE05,Jeong06,ESW07,Crocce08,Mat08}. Redshift distortions enhance a nonlinear damping along the line of sight direction\\citep[e.g.,][]{Meiksin99,SE05,ESW07,Mat08}.  Despite the resulting loss of the BAO signal with decreasing redshift, BAO are believed to be a robust standard ruler. The sound horizon scale is well determined from the CMB and the scale corresponds to $\\sim 100\\hMpc$ in present time, implying that the feature is still mostly on linear scales where evolutionary and observational effects are much simpler to predict than in the nonlinear regime.  Indeed, such degradation in the contrast of the BAO due to nonlinear structure growth, redshift distortions, and possibly galaxy bias has been studied since the late 90's and is relatively well understood \\citep[][]{Meiksin99,Springel05,Angulo05,SE05,White05,Eisen05,Jeong06,Crocce06b,ESW07,Huff07,Smith07,Matarrese07,Nishimichi07,Smith08,Angulo08,Crocce08,Mat08,Taka08,Sanchez08,Jeong09,Taruya09}.  Nonlinearity also induces a shift of the BAO scale in the low-redshift matter distribution relative to the sound horizon scale measured from the CMB \\citep[e.g., ][]{Smith08,Crocce08,Sanchez08}. As the demand for acoustic peak accuracy moves from the current level of a few \\% to the sub-percent level of future surveys, we are now required to understand the systematics on the BAO to much better precision. While it is evident that the shift of the BAO scale depends on the choice of the estimator, most recent studies seem to lay more weight on residual shifts using optimal estimators being at a sub-percent level at $z\\sim 0$ when accounting for nonlinear structure growth and redshift distortions \\citep{Crocce08,Sanchez08,SSEW08}, or even with halo/galaxy bias \\citep{Pad09}. Certainly more studies are necessary to confirm the results and ultimately reach a general consensus.   In the previous study \\citep{SSEW08} (hereafter SSEW08), we investigated effects of nonlinear evolution on BAO using a large volume (i.e., $320\\trihGpc$) of PM simulations. We found that the shift on the acoustic scale indeed increases with decreasing redshift and is less than a percent even at $z=0.3$ in redshift space.  In this work, we extend the study by using high-force-resolution simulations that are generated by a new \\Nb\\ code  ABACUS  by Metchnik \\& Pinto (in preparation). With the new simulations, we update the evolution of the shifts on the acoustic scale due to nonlinear structure growth and redshift distortions for various force, volume, and mass resolutions. We calculate propagators, i.e., the correlations between the linear density fields and the nonlinear density fields at low redshift, which directly manifest the nonlinear damping of the BAO. We also derive the mode-coupling contribution to the power spectrum with an attempt to qualitatively relate this to the measured values of the shifts. Recently, \\citet{PadLag09} showed that the density field after reconstruction is not the linear density field at second order. We relate the propagator and the mode-coupling term to an average and a dispersion of nonlinear density fields in the complex Fourier plane before and after reconstruction, relative to the linear density fields, and derive a signal-to-noise ratio of the standard ruler test from these two terms. The effect of galaxy bias will be presented in companion papers \\citep[Mehta et al. in preparation;][]{Xu09}.   It has been demonstrated that the original density-field reconstruction scheme based on the \\Zel\\ approximation, presented by \\citet{ESSS07}, is quite efficient for removing nonlinear degradation on BAO despite its simplicity \\citep{SE07,Huff07,SSEW08,PadLag09}. Efficiency of reconstruction in terms of an increase in the signal-to-noise depends on the redshift and the shot noise of the density fields. Meanwhile, it has been shown that the reconstruction removes almost all of the nonlinear shifts of the acoustic scale even when it seemingly is not efficient in terms of the signal-to-noise ratio (SSEW08).  This success, conversely, can be interpreted that a further improvement on the reconstruction scheme will be only a second order effect, at least for BAO. Nevertheless, we discuss  and test possible improvements on our fiducial scheme. While there are other sophisticated reconstruction methods in the literature aimed for the recovery of velocity fields and the initial density fields \\cite[e.g.,][]{MAK06},  in this paper we limit ourselves to mild modifications to our fiducial method, mainly due to its proven success for BAO.  We first test an implementation of 2LPT instead of the \\Zel\\ approximation \\citep{Zel70} and second, test iterative operations of the fiducial method in order to improve the signal-to-noise ratio.  The importance of an accurate prediction of the signal-to-noise ratio for future BAO surveys is evident. The calibration of the Fisher matrix-based estimations against the \\Nb\\ results have been tried repeatedly, and the resulting discrepancy is at most 20\\% (e.g., SSEW08). Further calibration is often limited by the volume of the simulations: in order to minimize {\\it the dispersion of the dispersion} among different measurements, we need a large number of random subsamples  while each subsample has an enough cosmic volume to measure the BAO scale.  In this paper, we  calibrate the Fisher matrix estimation based on \\citet{SE07} to a level of 1\\% by utilizing the enormous cosmic volume of $5000\\trihGpc$ from \\citet{Taka09}.    This paper is organized as following. In \\S~\\ref{sec:smethod}, we describe our new cosmological \\Nb\\ simulations and the methods of $\\chi^2$ analysis to measure the acoustic scales from the simulations. In \\S~\\ref{sec:salphas}, we present the resulting shifts and errors on the measurements of the acoustic scale when accounting for the nonlinear growth and redshift distortions before and after reconstruction. In \\S~\\ref{sec:spropagators} and \\S~\\ref{sec:Pmc}, we derive propagators and mode coupling terms from the simulations and qualitatively compare these with the errors and the mean values of the shifts from the simulations. In \\S~\\ref{sec:StoN}, we relate the propagator and the mode-coupling term to the signal-to-noise ratio of the standard ruler test. In \\S~\\ref{sec:stLPT}, we implement 2LPT and iteration into our reconstruction scheme and discuss the effect. In \\S~\\ref{sec:sRyuichi}, we use the $5000\\trihGpc$ of simulations by \\citet{Taka09} to test the Fisher matrix calculations in \\citet{SE07} given the minimal sample variance. Finally, in \\S~\\ref{sec:sdisc}, we summarize the major results presented in this paper. ",
        "conclusions": "\\label{sec:sdisc} We summarize the results presented in this paper.  1. We have measured shifts of the acoustic scale due to nonlinear growth and redshift distortions using three sets of simulations, G576, G1024, and T256. which differ in their force, volume, and mass resolution. The measured shifts from the various simulations are in agreement within $\\sim 1.5\\sigma$ of sample variance. We numerically find $\\al(z) -1 = (0.295\\pm 0.075)\\% [D(z)/D(0)]^{1.74\\pm 0.35}$ based on G576. If we fix the power index to be 2, as expected from the perturbation theory, we find the best fit of $\\al(z) -1 = (0.300\\pm 0.015)\\% [D(z)/D(0)]^{2}$.  2. We find a strong correlation with a non-unity slope between shifts in real space and redshift space. Meanwhile the correlation with the shifts at the initial redshift is weak. Reconstruction not only removes the mean shift and reduces errors on the mean, but also tightens these correlations. After reconstruction, we recover a slope of near unity for the correlation between the shifts in real and redshift space, and restore a strong correlation between the shifts at the low and the initial redshifts. We believe that, as the reconstruction removes shifts due to the second-order, nonlinear process in structure growth, the remaining shifts are dominated by the initial conditions.  3. We find that a propagator is well described by the \\Zel\\ approximation: for $z \\lesssim 1$, we find that the discrepancy, $\\Delta \\Signl$ is less than $0.5\\hMpc$. At high redshift, \\Zel\\ approximation seems to slightly underestimate the amount of nonlinear damping.  4. We have compared our measurements of shifts from $\\chi^2$ fitting with an oscillatory feature in $\\Pmc$ and find a qualitative agreement.  5. We construct the signal-to-noise ratio of the standard ruler test based on the measured propagator and mode-coupling term assuming a Gaussian error, and find a consistency with the measured errors on shift. We point out that the propagator describes the average projection of the nonlinear density field onto the linear density field while the mode-coupling term describes any random dispersion uncorrelated to the information of the linear density field. In the case of BAO, the recovery of BAO signal is well-described by $G(k)$, therefore, the average projection on the linear field, while the mode-coupling term contributes to the noise. On the other hand, the recovery of the broadband shape of the linear power spectrum will be hard to characterize with $G(k)$ alone, in the presence of $\\Pmc$.  6. We have attempted to improve our reconstruction method by implementing  2LPT and iterative operations. We find only a few \\% improvement in $G(k)$. We will further investigate variations in the 2LPT implementation in future work.  7. We test  Fisher matrix estimates of the uncertainty in the acoustic scale using $5000\\trihGpc$ of cosmological PM simulations (T256). At an expected sample variance level of 1\\%, the agreement between the Fisher matrix estimates based on \\citet{SE07} and the \\Nb\\ results is better than 10 \\%. \\\\  To conclude this paper, the acoustic peak shifts are small and can be accurately predicted, with control exceeding that required of the observational cosmic variance limit.  Moreover, reconstruction removes the shifts, decreases the scatter, and improves the detailed agreement with the initial density field, as hoped.  We have validated that the acoustic scale shift can be removed to better than 0.02\\% for the redshift-space matter field. We next plan to investigate the effects of galaxy bias (Mehta et al., in preparation). \\\\"
    },
    "0910/0910.0416_arXiv.txt": {
        "abstract": "Spectral data are presented for comets 2006~VZ13~(LINEAR), 2006~K4~(NEAT), 2006~OF2~(Broughton), 2P/Encke, and 93P/Lovas~I, obtained with the Cerro-Tololo Inter-American Observatory 1.5-m telescope in August 2007. Comet 2006~VZ13 is a new Oort cloud comet and shows strong lines of CN (3880~\\AA{}), the Swan band sequence for C$_2$ (4740, 5160, and 5630~\\AA{}), C$_3$ (4056~\\AA{}), and other faint species. Lines are also identified in the spectra of the other comets. Flux measurements of the CN, C$_2(\\Delta v=+1,0)$, and C$_3$ lines are recorded for each comet and production rates and ratios are derived. When considering the comets as a group, there is a correlation of C$_2$ and C$_3$ production with CN, but there is no conclusive evidence that the production rate ratios depend on heliocentric distance. The continuum is also measured, and the dust production and dust-to-gas ratios are calculated. There is a general trend, for the group of comets, between the dust-to-gas ratio and heliocentric distance, but it does not depend on dynamical age or class. Comet 2006~VZ13 is determined to be in the carbon-depleted (or Tempel~1 type) class. ",
        "introduction": "The icy nature of comets indicate they have been preserved at cold temperatures since the early stages of solar system formation. Consequently, they are commonly considered to be among the most primitive objects in the solar system. Determining their physical and chemical properties, and how they evolve, is important to our understanding of the formation of planetary systems, both our own and in general.  Several surveys have attempted to study and classify the chemical composition and evolution of comets. \\citet{A1995} conducted a photometric survey of 85 different comets over almost 20 years. Their findings indicated two major classes of comets: those that are carbon-depleted and those that are not. They found nearly all members of the carbon-depleted class are Jupiter family comets (JFCs), although not all JFCs are carbon-depleted. They also reported little variation of relative production rates with heliocentric distance or apparition; however, they noted a correlation between the dust-to-gas ratio and perihelion distance.  Three major spectroscopic surveys have also been conducted: \\citet{NS1989} reported spectrophotometry of 25 comets; \\citet{C1992} derived production rates for 17 faint comets; and \\citet{FH1996} surveyed the spectra of 39 comets into infrared wavelengths. The three surveys found slightly different results. \\citet{NS1989} reported a correlation between CN and dust, and that the C$_2$/CN production ratio changed with heliocentric distance. \\citet{C1992}, however, found the gas production ratios remained constant with activity level and heliocentric distance, except for NH$_2$/CN. \\citet{FH1996} concluded that most comets have roughly the same production rate ratios to within a factor of 2 or 3, although 10~per~cent of comets could be considered outliers.  More recently, \\citet{F2009} presented a spectroscopic survey of 92 comets over approximately 19 years. They report four taxanomic classes of comets: typical, Tempel~1 type, Giacobini-Zinner type, and the unusual object Yanaka (1988r). The typical comets have typical ratios of C$_2$, NH$_2$, and CN with respect to water, while Tempel-1 types have deficiencies in C$_2$ but normal NH$_2$ abundances. Giacobini-Zinner comets have low C$_2$ and NH$_2$ ratios with respect to water, while Yanaka has no detectable C$_2$ or CN emission, but normal NH$_2$ abundances. They conclude that the Halley family of comets (originating in the Saturn and Uranus region, but were scattered to the Oort cloud) shows no C$_2$ depletion, while objects originating in the Neptune region show a mixture of typical and C$_2$ depleted objects. Comets originating in the classical Kuiper belt form the C$_2$ depleted group.  Comet 2006~VZ13~(LINEAR) was discovered in November 2006 \\citep{S2006}, and is a new Oort cloud comet which passed perihelion on 10 August 2007. This presented a unique opportunity to observe a pristine comet as it passed perihelion. In addition, four other comets were observed with the same instrument and under the same observing conditions: 2006~K4~(NEAT), 2006~OF2~(Broughton), 93P/Lovas~I, and 2P/Encke. These comets represent a broad range of dynamical class, age, brightness and heliocentric distance. This allows for a comparison of production rates and ratios between comets of different classes and ages. \\begin{table*} \\centering \\begin{minipage}{150mm} \\caption{Orbital elements of the comets.\\setcounter{footnote}{\\value{footnote}}\\protect\\footnotemark}\\label{tab:orbits} \\begin{tabular}{@{}lcccccccc@{}} \\hline Comet & $a$ (au) & $e$ & $i$ ($^o$) & $q$ (au) & $Q$ (au) & $T_J$ & Perihelion Date (UT) & Comet Class\\setcounter{footnote}{\\value{footnote}}\\protect\\footnotemark \\\\ \\hline 2P/Encke\t&\t2.217\t&\t0.847\t&\t11.766\t&\t0.339\t&\t4.095\t& 3.026 & 2007 Apr 19 & NEO\\\\ 2006 K4\t\t&  1805.756  &\t0.998\t&\t111.333\t&\t3.189\t&\t3608.323 &\t-0.802 & 2007 Nov 29 & EOC\\\\ 2006 OF2\t&  -3367.756  &\t1.001\t&\t30.171\t&\t2.431\t&\tn/a \t& \tn/a & 2008 Sep 15 & NOC\\\\ 2006 VZ13\t&  -4083.355  &\t1.000\t&\t134.793 &\t1.015   &\tn/a\t&\tn/a & 2007 Aug 10 & NOC\\\\ 93P/Lovas~1\t&  4.391\t&\t0.612\t&\t12.218\t&\t1.705\t&\t7.078\t& 2.605 & 2007 Dec 17 & JFC\\\\ \\hline \\end{tabular} \\footnotetext{$^1$ All values are from the JPL Small-Body Database Browser; $a$ is the semi-major axis, $e$ is the eccentricity, $i$ is the inclination, $q$ is the perihelion distance, $Q$ is the aphelion distance, and $T_J$ is the Jupiter Tisserand parameter (given by $T_J=a_J/a +2\\cos i\\sqrt{a(1-e^2)/a_J}$).} \\footnotetext{$^2$ NEO = near earth object; EOC = evolved Oort cloud comet; NOC = new Oort cloud comet; JFC = Jupiter family comet.} \\end{minipage} \\end{table*}  \\setcounter{footnote}{0} \\begin{table*} \\centering \\begin{minipage}{105mm} \\caption{CTIO observational parameters.\\setcounter{footnote}{\\value{footnote}}\\protect\\footnotemark}\\label{tab:obs} \\begin{tabular}{@{}lccccc@{}} \\hline Comet & Date (UT) & $\\Delta$ (au) & $r_H$ (au) & $V$ (mag) & Exp. Time (s) \\\\ \\hline 2P\t\t&\t07/08/07 00:16\t&\t0.9400\t&\t1.9114\t&\t15.19\t&\t3600\\\\ &\t07/08/07 23:56\t&\t0.9550\t&\t1.9221\t&\t15.30\t&\t3650\\\\ &\t13/08/07 01:30\t&\t1.0381\t&\t1.9781\t&\t15.92\t&\t5400\\\\ &\t14/08/07 01:14\t&\t1.0551\t&\t1.9890\t&\t16.04\t&\t5400\\\\ 2006 K4\t\t&\t07/08/07 01:24\t&\t2.6270\t&\t3.3713\t&\t15.75\t&\t3600\\\\ &\t08/08/07 02:08\t&\t2.6364\t&\t3.3681\t&\t15.75\t&\t3600\\\\ &\t13/08/07 03:33\t&\t2.6869\t&\t3.3529\t&\t15.78\t&\t5400\\\\ &\t14/08/07 03:04\t&\t2.6977\t&\t3.3500\t&\t15.78\t&\t5400\\\\ &\t15/08/07 02:33\t&\t2.7084\t&\t3.3472\t&\t15.79\t&\t5400\\\\ 2006 OF2\t&\t07/08/07 02:53\t&\t3.7934\t&\t4.7860\t&\t16.33\t&\t3600\\\\ &\t08/08/07 03:19\t&\t3.7825\t&\t4.7782\t&\t16.32\t&\t3600\\\\ &\t13/08/07 06:32\t&\t3.7316\t&\t4.7382\t&\t16.26\t&\t5400\\\\ &\t14/08/07 06:17\t&\t3.7226\t&\t4.7304\t&\t16.25\t&\t5400\\\\ &\t15/08/07 05:06\t&\t3.7143\t&\t4.7229\t&\t16.24\t&\t5400\\\\ 2006 VZ13\t&\t06/08/07 23:14\t&\t1.0207\t&\t1.0175\t&\t13.26\t&\t1800\\\\ &\t07/08/07 23:14\t&\t1.0482\t&\t1.0165\t&\t13.32\t&\t1800\\\\ &\t12/08/07 23:22\t&\t1.1866\t&\t1.0159\t&\t13.58\t&\t3600\\\\ &\t13/08/07 23:19\t&\t1.2142\t&\t1.0166\t&\t13.64\t&\t3600\\\\ 93P\t\t&\t08/08/07 06:03\t&\t1.4603\t&\t2.1451\t&\t16.61\t&\t7200\\\\ &\t15/08/07 07:36\t&\t1.3617\t&\t2.1059\t&\t16.38\t&\t9600\\\\ \\hline \\end{tabular} \\footnotetext{$^1$ The date indicates the start time of the first observation of the comet; $\\Delta$ is the geocentric distance, $r_H$ is the heliocentric distance, and $V$ is the total magnitude. The exposure time is the total of all observations on the given night.} \\end{minipage} \\end{table*} In this paper, spectroscopic observations are presented of five comets, obtained in August 2007. The production rates of CN, C$_2$, and C$_3$ are calculated, and the production ratios with respect to CN are derived. In addition, dust production rates and dust-to-gas ratios are calculated. ",
        "conclusions": "Spectroscopic observations for five comets are reported. From these spectra, the gas production rates and production rate ratios have been calculated. There seems to be a linearity of the production rate ratios with respect to CN, agreeing with past studies \\citep{C1987,A1995}. There does not seem to be a correlation between the production rate ratios and heliocentric distance. By comparing the spectra of comet 2006~VZ13 with those presented in \\citet{F2009}, it is determined to be a carbon-depleted (Tempel~1 type) comet.  The dust production rate and the dust-to-gas ratio are also calculated for each comet. The ratio stays relatively constant for each comet, except for 2P due to a variability in CN production. There seems to be an overall dependence on heliocentric distance when considering the comets as a group. There is no observed dependence of the dust-to-gas ratio on the dynamical age or class of the comets."
    },
    "0910/0910.5719_arXiv.txt": {
        "abstract": "White dwarf--neutron star binaries generate detectable gravitational radiation. We construct Newtonian equilibrium models of corotational white dwarf--neutron star (WDNS) binaries in circular orbit and find that these models terminate at the Roche limit.  At this point the binary will undergo either stable mass transfer (SMT) and evolve on a secular timescale, or unstable mass transfer (UMT), which results in the tidal disruption of the WD. The path a given binary will follow depends primarily on its mass ratio. We  analyze  the  fate of known WDNS binaries and use population synthesis results to estimate the number of LISA-resolved galactic binaries that will undergo either SMT or UMT. We model the quasistationary SMT epoch by solving a set of simple ordinary differential equations and compute the corresponding gravitational waveforms. Finally, we discuss in general terms the possible fate of binaries that undergo UMT and construct approximate Newtonian equilibrium configurations of merged WDNS remnants. We use these configurations to assess plausible outcomes of our future, fully relativistic simulations of these systems. If sufficient WD debris lands on the NS, the remnant may collapse, whereby the gravitational waves from the inspiral, merger, and collapse phases will sweep from LISA through LIGO frequency bands. If the debris forms a disk about the NS, it may fragment and form planets. ",
        "introduction": "The inspiral and merger of compact binaries represent some of the most promising sources of gravitational waves (GWs) for detection by ground-based laser interferometers like LIGO \\cite{LIGO1,LIGO2}, VIRGO \\cite{VIRGO1,VIRGO2}, GEO \\cite{GEO}, TAMA \\cite{TAMA1,TAMA2} and AIGO \\cite{AIGO}, as well as by proposed space-based interferometers like LISA \\cite{LISA} and DECIGO \\cite{DECIGO}. Extracting physical information from gravitational radiation emitted by compact binaries requires careful modeling of these systems (see \\cite{BSReview} for a review). Most effort to date has focused on modeling black hole--black hole (BHBH) binaries \\cite{Pretorius2005a,Baker2006,Campanelli2006,Brugmann2006a,Herrmann2006,Sperhake2006,Etienne2007a,Healy,Baker2006b,Lousto2007,Gonzalez08,Scheel2008} and neutron star--neutron star (NSNS) binaries \\cite{2004PhRvD..69l4036F,2005A&A...431..297B,2005PhRvD..71h4021S,2006PhRvD..73f4027S,2008PhRvD..77b4006A,2008PhRvD..78b4012L,2008PhRvD..78h4033B}, with some recent work on black hole--neutron star (BHNS) binaries \\cite{Rantsiou08, Loffler06, Faber06, Shibata07, Shibata08,Yamamoto08, Etienne08a, Etienne08, Duez08}.  In this paper we begin to explore WDNS binaries. They are plausible sources of low-frequency GWs for LISA, DECIGO and possibly also, as we shall see, high-frequency GWs for LIGO, VIRGO, GEO, TAMA and AIGO.  Like NSNS binaries, WDNS binaries are known to exist. We have compiled a list of observed WDNS binaries % in Tables \\ref{table1} and \\ref{table2}. Table~\\ref{table1} lists those binaries with relatively well-determined individual component masses. By ``relatively well-determined\" we mean that in determining the masses no assumption was made about the mass of the NS. Table~\\ref{table2} lists those binaries for which the masses of the individual components have not been well-determined as yet. For these cases the NSs have been assumed to have a mass of either $1.35 M_\\odot$ or $1.40 M_\\odot$. Of all the objects presented in Tables \\ref{table1} and \\ref{table2} only B2303+46, J1141-6545, J0751+1807 and J1757-5322 have been positively identified as a pulsar with a WD companion. For all other systems there is at best strong evidence that the companion of the observed pulsar is a WD.  The frequency of GWs emitted by the binary with the shortest period, J1141$-$6545, is $\\simeq 1.2\\times 10^{-4} \\rm Hz$. This frequency lies in the expected band of LISA, but well below the cutoff  frequency of $\\simeq 10^{-3} \\rm Hz$ due to the double WD confusion background \\cite{Nelemans01}.  It would therefore be impossible to resolve this signal.    Assuming the quadrupole approximation for a WDNS binary with masses $M_{\\rm WD}$, $M_{\\rm NS}$, reduced mass $\\mu$, and total mass $M_{\\rm T}=M_{\\rm NS}+M_{\\rm WD}$ in circular orbit with angular frequency $\\Omega$, the amplitude of GWs coming from this object is $h = 4 \\mu M_{\\rm T} / (r A)$ \\cite{Shapiro}, where $r$ is the distance to the object and $A$ the binary separation. In \\cite{distanceJ11} a lower bound is given to the distance to PSR J1141$-$6545 of $r\\geq 3.7$ kpc. Using the data of Table \\ref{table1} we find that the maximum amplitude of the expected waves coming from this object is $h_{max}=1.36\\times 10^{-23}$. A gravitational wave of this amplitude is below LISA's sensitivity at $10^{-4} \\rm Hz$ ($h_s\\sim 10^{-21}$ ) and hence undetectable. In order for detectable gravitational wave signals to be emitted the orbital separation of PSR J1141$-$6545 has to decrease by a factor of about 100.   The emission of gravitational radiation will cause a binary to inspiral to a close, nearly circular orbit.  In this regime the binary can be described by a sequence of quasiequilibrium configurations.  Once the separation of a WDNS binary becomes a few times larger than the size of the WD, the amplitude of the emitted GWs, as we will see,  is large enough to allow for detection by space-based gravitational wave observatories. An important issue to address then is the number of Galactic WDNS binaries that could be detected by space-based interferometers.   \\begin{table*}[t] \\caption{WDNS binaries with well-determined masses. From left to right the columns give the name of the object, the orbital period, the corresponding quadrupole GW frequency, the WD mass, the NS mass, the total mass, the mass ratio $q=M_{\\rm WD}/M_{\\rm NS}$ and the stability of mass transfer after its onset.} \\begin{tabular}{lccccccc}\\hline\\hline \\multicolumn{1}{p{2.15 cm}}{\\hspace{0.3 cm} PSR } & \\multicolumn{1}{p{2.cm}}{\\hspace{0.25 cm} $T$ (days)} & \\multicolumn{1}{p{3.5cm}}{\\hspace{0.5 cm} $f_{\\rm GW}=\\frac{2}{T}$ $(10^{-4} \\rm Hz)$ \\quad} & \\multicolumn{1}{p{2cm}}{ \\hspace{0.35 cm}$M_{\\rm WD}(M_{\\odot})$ \\qquad} & \\multicolumn{1}{p{2cm}}{\\hspace{0.3 cm} $M_{\\rm NS}(M_{\\odot})$ \\qquad }  & \\multicolumn{1}{p{2cm}}{\\hspace{0.3 cm} $M_{\\rm T}(M_{\\odot})$ \\qquad }  & \\multicolumn{1}{p{1.5cm}}{\\hspace{0.5 cm} $q$ \\quad} & \\multicolumn{1}{p{1.5cm}}{\\hspace{0.15 cm} Stable? \\quad} \\vspace{0.05cm} \\\\ \\hline B2303$+$46\\footnotemark[1]   \t& $12.34$  &  $0.0187$                   & $1.3$                 &   $ 1.34$     &   $2.64$     & $0.97$       & No                            \\\\  \\hline J0621$+$1002\\footnotemark[2] \t& $8.32$    &  $0.0278$                   & $0.67$               & $ 1.70$      &   $2.37$      & $0.394$     & ?                             \\\\  \\hline J1141$-$6545\\footnotemark[3]$^,$\\footnotemark[4]& $0.198$  &  $1.169$                     & $1.02$               &  $ 1.27$     &   $2.29$      & $0.803$     & No     \\\\  \\hline B1516$+$02B\\footnotemark[5]    \t& $6.858$  &  $0.0337$                   & $0.13$               & $ 2.08$      &   $2.21$      & $0.0625$    & Yes                             \\\\  \\hline J1713$+$0747\\footnotemark[1]    \t& $67.8$    &  $0.0034$                   & $0.33$               &  $ 1.60$     &   $1.93$      & $0.206$     & ?                       \\\\  \\hline B1855$+$09\\footnotemark[1]   \t       & $12.3$    &  $0.0188$                   & $0.267$             &  $ 1.58$     &   $1.847$    & $0.169$     & Yes                       \\\\  \\hline J0437$-$4715\\footnotemark[1]  \t& $5.74$    &  $0.0403$                   & $0.236$             &  $ 1.58$     &   $1.816$    & $0.149$     & Yes                        \\\\  \\hline J1012$+$5307\\footnotemark[1]  \t& $0.605$  &  $0.382$                     & $0.16$                & $ 1.64$     &   $1.80$      & $0.097$     & Yes                       \\\\  \\hline J0751$+$1807\\footnotemark[1]  \t& $0.263$  &  $0.88$                       & $0.125$              & $ 1.26$     &   $1.385$    & $0.099$     & Yes                       \\\\  \\hline\\hline \\end{tabular} \\begin{flushleft} \\footnotetext[1]{Stairs \\cite{WDNS7} and references therein.}  % \\footnotetext[2]{Nice et al. \\cite{WDNS2}.}   % \\footnotetext[3]{Bhat et al. \\cite{WDNS4}.}  % \\footnotetext[4]{Bailes et al. \\cite{WDNS5}.}   % \\footnotetext[5]{Freire et al. \\cite{WDNS6}.}   % \\end{flushleft} \\label{table1} \\end{table*}   \\begin{table*} [t] \\caption{WDNS binaries which do not have well-determined masses. From left to right the columns give the name of the object, the orbital period, the corresponding quadrupole GW frequency, the WD mass, the NS mass, the total mass, the mass ratio $q=M_{\\rm WD}/M_{\\rm NS}$ and the stability of mass transfer after its onset.} \\begin{tabular}{lccccccc}\\hline\\hline \\multicolumn{1}{p{2.15 cm}}{\\hspace{0.3 cm} PSR } & \\multicolumn{1}{p{2.cm}}{\\hspace{0.25 cm} $T$ (days)} & \\multicolumn{1}{p{3.5cm}}{\\hspace{0.5 cm} $f_{\\rm GW}=\\frac{2}{T}$ $(10^{-4} \\rm Hz)$ \\quad} & \\multicolumn{1}{p{2cm}}{ \\hspace{0.35 cm}$M_{\\rm WD}(M_{\\odot})$ \\qquad} & \\multicolumn{1}{p{2cm}}{\\hspace{0.3 cm} $M_{\\rm NS}(M_{\\odot})$ \\qquad }  & \\multicolumn{1}{p{2cm}}{\\hspace{0.3 cm} $M_{\\rm T}(M_{\\odot})$ \\qquad }  & \\multicolumn{1}{p{1.5cm}}{\\hspace{0.5 cm} $q$ \\quad} & \\multicolumn{1}{p{1.5cm}}{\\hspace{0.15 cm} Stable? \\quad} \\vspace{0.05cm} \\\\ \\hline J1435$-$60\\footnotemark[1]   \t& $1.355$  &  $0.1708$                   & $1.10$                 & $ 1.40$      &   $2.50$      & $0.785$         & No                             \\\\  \\hline J1157$-$5114\\footnotemark[2]   \t& $3.507$  &  $0.066$                     & $1.14$                 &  $ 1.35$     &   $2.49$      & $0.844$         & No                             \\\\ \\hline J1453$-$58\\footnotemark[1]     \t& $12.42$  &  $0.0186$                   & $1.07$                 &   $ 1.40$     &   $2.47$     & $0.764$       & No                            \\\\  \\hline J1022$+$1001\\footnotemark[1]    & $7.805$  &  $0.0296$                   & $0.872$               & $ 1.40$    &   $2.272$     & $0.623$       & No                             \\\\  \\hline B0655$+$64\\footnotemark[1]    \t& $1.029$   &  $0.2948$                   & $0.814$               &  $ 1.40$     &   $2.214$      & $0.581$         & No                       \\\\  \\hline J2145$-$0750\\footnotemark[1]   \t& $6.839$    &  $0.0338$                   & $0.515$               &  $1.40$     &   $1.915$    & $0.368$         & ?                       \\\\  \\hline J1757$-$5322\\footnotemark[2]  \t& $0.453$    &  $0.511$                   & $0.55$              &  $1.35$     &   $1.90$    & $0.407$         & ?                        \\\\  \\hline J1603$-$7202\\footnotemark[1]  \t& $6.309$    &  $0.0367$                   & $0.346$              &  $ 1.40$     &   $1.746$    & $0.247$         & ?                        \\\\  \\hline J1810$-$2005\\footnotemark[1]    & $15.01$    &  $0.0154$                   & $0.34$              &  $ 1.40$     &   $1.74$    & $0.243$         & ?                        \\\\  \\hline J1904$+$04\\footnotemark[1]  \t& $15.75$  &  $0.0147$                     & $0.27$                & $ 1.40$     &   $1.67$      & $0.193$         & ?                       \\\\  \\hline J1232$-$6501\\footnotemark[1]    & $3.507$  &  $0.066$                       & $0.175$              & $ 1.40$     &   $1.575$    & $0.125$         & Yes                       \\\\  \\hline\\hline \\end{tabular} \\begin{flushleft} \\footnotetext[1]{Tauris et al. \\cite{WDNS1} and references therein.}  % \\footnotetext[2]{Edwards and Bailes \\cite{WDNS3}.}   % \\end{flushleft} \\label{table2} \\end{table*}   Population synthesis calculations by Nelemans et al.\\ \\cite{Nelemans01} show that there are about $2.2\\times 10^{6}$ WDNS binaries in our Galaxy, and that they have a merger rate of $1.4\\times 10^{-4} \\rm yr^{-1}$. Furthermore, Nelemans et al.\\ find that after a year of integration, LISA should be able to detect $128$ WDNS binaries and, after considering the contribution of the double WD background GW noise, resolve $38$ of these. On the other hand, calculations by Cooray \\cite{Cooray2004}, give much more conservative numbers of resolved WDNS binaries. In particular, Cooray finds that the number of LISA-resolved WDNS binaries ranges between $1$--$10$, using a WDNS merger rate between $10^{-6} \\rm yr^{-1}$--$10^{-5} \\rm yr^{-1}$. Cooray's upper limit was based on merger rates calculated by Kim et al.\\ \\cite{Kim2004}.  In this work we focus on WDNS binaries in close binary separations, and examine the termination point of quasiequilibrium sequences describing such binaries.  Several different astrophysical scenarios can result in such a termination point for binaries in general, including {\\em direct contact} of the binary components, {\\em Roche lobe overflow} by one of the two companions, or the binary reaching an {\\em innermost stable circular orbit} (ISCO).  As we will find, quasiequilibrium sequences of WDNS terminate when the WD fills its Roche lobe, resulting in mass transfer from the WD onto the NS across  the inner Lagrange point.  This mass transfer can either be stable (SMT) or unstable (UMT).  We also refer to the latter scenario as the tidal disruption of the WD by the NS.  To determine which of these two outcomes is likely for a given system, we will follow the approach of Verbunt and Rappaport \\cite{Verbunt88} and Faber et al. \\cite{Faber}. As indicated in our Tables~\\ref{table1} and~\\ref{table2}, we shall find that among the observed WDNS binaries, some will undergo SMT, while others will undergo tidal disruption (i.e.~UMT). Note that mass transfer stability has also been studied in \\cite{Hut,Marsh}.   In the case of SMT, the orbital evolution of the binary occurs on a secular timescale determined by the emission of gravitational radiation. Therefore, the quasistationary conservative treatment of Clark and Eardley \\cite{Clark77} and Faber et al. \\cite{Faber} is adequate to follow the evolution during this secular phase. Note that a quasistationary treatment to follow the orbital evolution of binary systems has been employed by several authors in the past. For example, Rappaport et al.~\\cite{Rappaport82} studied compact binaries where the mass  of the secondary can be up to $\\sim 1M_\\odot$ and modeled as a $n=3/2$ polytrope. Fryer et al.~\\cite{Fryer99} employed a non-conservative quasistationary approach to study the evolution of white dwarf--black hole binaries, while Marsh et al.~\\cite{Marsh} studied the evolution of double WD binary systems. Both the tidal disruption and SMT phase of double WD systems have also been studied in \\cite{Benz1990, Podsiadlowski92, RasioShapiro95, Segretain1997, Guerrero2004, Yoon2007,Dan08} via SPH simulations and in \\cite{Dsouza,Motl} via grid-based hydrodynamic calculations, all in Newtonian gravitation. Finally, Newtonian SPH simulations of encounters of WDs with intermediate mass BHs ($M_{\\rm BH}\\sim 10^3 M_\\odot$) have been performed in \\cite{RosswogWDBH1,RosswogWDBH2}.  In the case of tidal disruption, on the other hand, the system will evolve on a hydrodynamical (orbital) timescale. In this scenario the NS may plunge into the WD and spiral toward the center of the star liberating its gravitational potential energy as heat in the WD material. Alternatively, the NS may be the receptacle of massive debris from the disrupted WD. Depending on the details of the equation of state, a cold degenerate gas can support a maximum NS rest-mass between $1.89 M_\\odot-2.67 M_\\odot$ (corresponding to  a gravitational mass between $1.65M_\\odot-2.20M_\\odot$) against catastrophic collapse if it is not rotating (the OV limit), about $20\\%$ more mass if it is rotating uniformly  (a ``supramassive NS'', e.g.~\\cite{CooST94}), and at least $50\\%$ more mass if it rotates differentially (a ``hypermassive NS'') \\cite{BauSS00,2004ApJ...610..941M}.  The fate of the merged WDNS then depends on the initial masses of the progenitor stars, the degree of mass and angular momentum loss during the WD disruption and binary merger phases, the angular momentum profile of the NS remnant and the extent to which the remnant gas is heated by shocks as it pours onto the NS and forms an extended, massive mantle. These are issues that require a hydrodynamic simulation to resolve.  Moreover, ascertaining whether or not the neutron star ultimately undergoes a catastrophic collapse to a black hole (either prompt or delayed) requires that such a simulation be performed in full general relativity.  We plan to explore some of these alternative hydrodynamical scenarios in detail in the future,  aided by simulations that employ our adaptive mesh refinement (AMR) relativistic hydrodynamics code \\cite{Etienne08}.  In this paper we survey the problem in qualitative terms. In Section~\\ref{Sec:Eq} we identify some of physical parameters that are likely to play a key role in the evolution of a WDNS binary system. We then model the secular inspiral epoch of the binary by constructing Newtonian equilibrium models of corotational WDNS binaries in close circular orbit, up to the Roche limit.  In Section~\\ref{MTSvsTD} we follow the stability analysis of Verbunt and Rappaport \\cite{Verbunt88} in order to determine the late-evolution of these binaries (SMT vs.~tidal disruption). We treat the quasistationary SMT epoch by applying the approach of Clark and Eardley \\cite{Clark77} and Faber et al.~\\cite{Faber} and compute the corresponding gravitational waveforms. We discuss in general terms the possible outcomes of binaries that undergo UMT and construct approximate equilibrium configurations of merged WDNS remnants. We use these models to make some predictions regarding our future fully relativistic simulations of these systems. Finally, we conclude in Section \\ref{summary}, where we summarize the main findings of this work. ",
        "conclusions": "\\label{summary}  In this paper we considered WDNS binaries and ``set the stage\" for fully relativistic hydrodynamic simulations of the WDNS merger. Like NSNS binaries, but unlike BHNS and BHBH binaries, WDNS binaries are known to exist and are abundant. We have compiled a list of observed WDNS binaries and their properties in Tables \\ref{table1} and \\ref{table2}.  The emission of gravitational radiation will cause the orbital separation of a WDNS binary to shrink. We modeled corotational WDNS binaries in circular orbit and found that these models terminate at the Roche limit. At this point the binary can undergo either SMT and evolve on a secular timescale or UMT and evolve on a hydrodynamical timescale.  Following the stability analysis of Verbunt and Rappaport \\cite{Verbunt}, we showed that  the subsequent fate of the binary is determined by a critical mass ratio. Using the results of this stability analysis we predicted the possible fates of known WDNS binaries and we indicated this in our Tables~\\ref{table1} and~\\ref{table2}. Furthermore, based on population synthesis results by Nelemans et al. \\cite{Nelemans01}, we gave an estimate of the number of LISA-resolved galactic binaries per year that will undergo SMT and the number of those that will undergo tidal disruption. We found that approximately up to $16$ WDNS binaries will undergo UMT (tidal disruption), and up to $20$ SMT.  In the case of SMT, the timescale of the mass transfer, and hence the timescale of the binary evolution, is set by the emission of gravitational radiation. We treated this quasistationary SMT epoch by applying the approach of Clark and Eardley \\cite{Clark77} and Faber et al. \\cite{Faber}, also adopted in \\cite{Rappaport82, Fryer99,Marsh}, and we estimated the corresponding gravitational waveforms.  We also constructed approximate equilibrium configurations of rigidly rotating WDNS merged remnants in order to explore possible outcomes of WDNS mergers.  We found that unless some process removes angular momentum from the system or leads to the formation of a massive disk, these massive equilibrium configurations cannot collapse directly to form a Kerr black hole since their angular momentum ($J/M^2\\sim 20>1$) exceeds the Kerr limit. Furthermore, in most of our case studies we found that the merged remnants have a ratio of spin kinetic to gravitational potential energy  greater than $0.27$, which implies that they are dynamically unstable to bar formation.  However, we emphasize that these preliminary results are at best approximate and must be confirmed and/or revised by detailed numerical simulations.  The fate of a merged WDNS binary depends on the initial mass of the cold progenitor stars, the degree of mass and angular momentum loss during the WD disruption and binary merger phases, the angular momentum profile of the WDNS remnant and the extent to which the remnant gas is heated by shocks as it pours onto the NS and forms an extended, massive mantle. These issues all require a hydrodynamic simulation to resolve.  Moreover, ascertaining whether or not the NS ultimately undergoes a catastrophic collapse to a BH (either prompt or delayed) requires that such a simulation be performed in full general relativity.  In fact, even the final fate of the NS in the alternative scenario in which there is a long epoch of SMT may also lead to catastrophic collapse, if the final NS mass exceeds the TOV mass limit. This scenario too will require a general relativistic hydrodynamic simulation to track. We plan to explore some of these alternative scenarios in detail in the future, aided by simulations that employ our AMR relativistic hydrodynamics code."
    },
    "0910/0910.5233_arXiv.txt": {
        "abstract": "We calculate the distribution of neutral Hydrogen within 750 proper $\\chikps$ of a quasar, $L_{\\rm bol} = 1.62 \\times 10^{13} {\\rm L}_{\\sun} \\approx L_{\\rm Edd}$, powered by accretion onto a super massive black hole, $M_{\\rm BH} = 4.47 \\times 10^8 {\\rm M}_{\\sun}$, at z = 3.  Our numerical model includes cosmological initial conditions, gas dynamics, star formation, supernovae feedback, and the self consistent growth and thermal feedback of black holes calculated using \\Gadget as well as a detailed post-processing ray tracing treatment of the non-uniform ionizing radiation field calculated using \\Sphray.  Our radiative transfer scheme naturally accounts for the self shielding of optically thick systems near the luminous central source. We show that the correct treatment of self shielding introduces a flattening feature into the neutral column density distribution around Log $N_{\\rm HI}$ = 20 and that regions with the lowest neutral fractions are not necessarily those with the highest density gas.  For comparison with our numerical work, we solve a Ricatti equation which determines the equilibrium Hydrogen ionization fractions in the presence of a radiation field that falls off as one over $r$ squared with regions above a given gas density threshold completely shielded from ionizing radiation. We demonstrate that these simple semi analytic models cannot reproduce the neutral Hydrogen field calculated using {\\small SPHRAY}. We conclude by comparing our models of this single proximity zone to observations by Hennawi and Prochaska of the absorption spectra of background quasars which are coincident on the sky with foreground quasars in their Quasars Probing Quasars (QPQ) series of papers.  Compared to the QPQ sample, we find a factor of 3 fewer optically thick (Log $N_{\\rm HI} \\ge 17.2$) systems around our quasar, however the dark matter halo that hosts our simulated quasar, $M_{\\rm Halo} = 5.25 \\times 10^{12} {\\rm M}_{\\sun}$, is less massive than the typical QPQ host halo by a factor of four. Allowing for a linear scaling between halo mass, baryonic overdensity and number of absorbers, we estimate the typical host halo mass in the QPQ sample as $1.92 \\times 10^{13} {\\rm M}_{\\sun}$. ",
        "introduction": "After the reionization of the Universe by the first luminous objects, the majority of neutral Hydrogen resides in gravitationally collapsed objects, specifically Damped Lyman-$\\alpha$ systems (DLAs). These systems are observed via absorption lines in the spectra of distant quasars and are historically identified as those absorbers with HI column densities $N_{\\rm HI} > 2 \\times 10^{20} {\\rm cm^{-2}}$ with lower column density systems $10^{17.2} < N_{\\rm HI} < 2 \\times 10^{20} {\\rm cm^{-2}}$ being labeled Lyman Limit Systems \\cite[LLSs, see][for a review]{2005ARA&A..43..861W}. In contrast, the systems that give rise to the Lyman-$\\alpha$ forest have $N_{\\rm HI} < 10^{17.2} {\\rm cm^{-2}}$ \\cite[see][for reviews]{1998ARA&A..36..267R, 2003AIPC..666..157W}. This historical column density threshold ($N_{\\rm HI} = 10^{17.2} {\\rm cm^{-2}}$) serves to divide systems into those that are predominantly neutral (the DLAs and LLSs) and those that reside in the mostly ionized intergalactic medium (IGM).  The DLAs and LLSs remain mostly neutral in the presence of an ionizing ultraviolet (UV) background and local point sources through self shielding.  Studying them in absorption opens a window into the post reionization population of cold, dense, neutral gas and provides a survey technique with a bias complimentary to that of emission surveys.  To study these systems numerically requires treating the UV ionizing background at some level.  Many cosmological simulations have followed in the steps of \\cite{HM96} and \\cite{1996ApJS..105...19K} by considering a spatially uniform, time variable, background with a spectral shape characteristic of quasar and stellar radiation that has been reprocessed by the IGM.  The background plays a role in determining the cooling function of cosmological gas and so its inclusion is also necessary at some level for a realistic description of galaxy and star formation.  A uniform background is a good first approximation if one is interested in radiative cooling, however it ignores completely the self shielding that defines the DLAs and LLSs.  The most straight forward way to account for this shielding is to reduce the UV field to zero in regions above a given gas density threshold as was done in \\cite{1998ApJ...495..647H}.  A more detailed treatment based on the solution for plane parallel radiation incident on a constant density slab is described in \\cite{KWHM96} and used in the work of \\cite{1997ApJ...484...31G, 2001ApJ...559..131G, 1997ApJ...486...42G}. These corrections are not based on transferring radiation through the simulation volume, but rather on applying the plane parallel solution on a pixel by pixel basis to HI column density maps in post processing.  Other approaches to self shielding include that of \\cite{2003ApJ...598..741C} who include attenuation of the UV background on a cell by cell basis akin to \\cite{KWHM96}. \\cite{2006ApJ...645...55R}, who include a treatment of the transfer of ionizing radiation in post processing using the Fully Threaded Transport Engine ({\\small FTTE}) described in \\cite{2005MNRAS.362.1413R}.  And \\cite{2007ApJ...660..945N}, who use a multi phase gas model to treat star formation and assume the cold dense phase to be fully neutral.  Recently, a series of nested galaxy formation simulations at different resolutions was used by \\cite{2008MNRAS.390.1349P} to study DLAs over a broad range of mass scales.  They include a simplified radiative transfer scheme to account for self shielding which lies somewhere between the full radiative transfer modeling and the pixel by pixel corrections.  The works mentioned above were aimed at studying systems where the radiation field is not dominated by a local point source. \\cite{2005ApJ...620L..91M} applied simple analytic arguments based on the conservation of surface brightness to argue that {\\it on average} the effect from local sources is negligible compared to that of the background for systems with optical depths below that of Lyman Limit Systems.  On the other hand, \\cite{2006ApJ...643...59S} use analytic arguments to show the local radiation field is likely to be important for denser systems that tend to cluster around the large scale overdensities that host these sources.  In this work, we examine the balance between the UV background, local sources, and self shielding by combining, for the first time, a hydrodynamic simulation that tracks the formation and accretion history of black holes with a detailed ray tracing treatment of the non-uniform UV radiation field.  Shielding is most important in the presence of dense gas and strong UV fields, both of which occur near active galactic nuclei (AGN) and quasars.  In fact, the transition between the background UV field and the local AGN/quasar UV field serves to define the proximity region of these objects.  With the availability of large quasar catalogues and high resolution spectroscopy, it has become possible to study absorbers proximate to a foreground quasar in the spectrum of a coincident background quasar.  Work such as this has been carried out in a series of papers entitled ``Quasars Probing Quasars'', the first of which is authored by \\cite{QpQI} (HP06 in the rest of this paper).  We compare our numerical models to observations from this body of work and to theoretical calculations of the HI column density by \\cite{2002ApJ...568L..71Z}.   The format of this paper is as follows.  In \\S2 we describe the \\Gadget simulation that determines our temperature and density fields, in \\S3 we review the ray tracing code \\Sphray used to calculate the transfer of ionizing radiation through this density field, in \\S4 we describe how we model our sources of ionizing radiation, in \\S5 we describe our semi analytic model, in \\S6 we describe our ray tracing results and compare them to the semi analytic model, in \\S7 we compare our theoretical results with those of \\cite{QpQI}, and \\cite{2002ApJ...568L..71Z} and in \\S8 we conclude and discuss further work. ",
        "conclusions": "In this paper, we have presented a detailed model for the neutral hydrogen gas in the proximity zone of a quasar powered by accretion onto a super massive black hole.  To do this, we combined a cosmological hydrodynamic \\Gadget simulation of the formation and growth of black holes with a detailed radiative transfer post processing code {\\small SPHRAY}. We focused our attention on self shielding of the neutral gas and, by construction, the contribution from a local central source.  Much of the previous work on Lyman Limit and Damped Lyman $\\alpha$ systems ignored the contribution from point sources assuming that the background field dominates.  This is true, on average, for absorbers less dense then Lyman Limit systems but is not true for the more dense absorbers which cluster around luminous point sources and those underdense systems that just happen to be near sources.  In addition it is crucial that these sorts of models be constructed to interpret quasar pair observations such as those in HP06.  We find that the results of our ray tracing algorithm \\Sphray cannot be reproduced with simple semi analytic models in which gas above a given density threshold is completely shielded from ionizing radiation.  This is because the photoionization rate at any given point is determined by the partial shielding provided by the local density field and to a lesser degree the partial shadowing seen in Figure \\ref{GHIhealpix}.  We have shown in Figure \\ref{NHIpdf} that our \\Gadget + \\Sphray models naturally reproduce the observed shape of the $N_{\\rm HI}$ column density distribution.  We have also shown that the flattening feature around Log $N_{\\rm HI} = 20$ attributed to self shielding in the analytic work of \\cite{2002ApJ...568L..71Z} is not due entirely to self shielding, but is largely enhanced by it.   When compared to the observations of HP06, our ray tracing and semi analytic models fall short of capturing the number of dense absorbers proximate to luminous quasars.  We have described improvements to our model that would bring our results into agreement with these observations including sampling more massive host halos, correcting for the under resolved clumping of gas, allowing obscuration / periodic emission of the quasar radiation and using a larger simulation volume to better model uncertain radial separations.  These issues are not intractable and can be dealt with using a larger low resolution simulation to sample more massive dark matter host halos coupled with a high resolution resimulation of a more representative proximity zone.   Finally we have shown that the total amount of neutral HI within the 750 ${h^{-1}}$ kpc region around our central quasar can change by up to an order of magnitude depending on the shielding prescription used.  This means that HI surveys that do not resolve individual galaxy features but whose goal is to measure the integrated signal from a proximity zone will need to take these effects into account.  This is on scales much smaller then the intensity mapping proposed by \\cite{2009arXiv0902.3091P} to measure Baryon Acoustic Oscillations, but would be relevant for connecting HI or 21 cm galaxy surveys to the underlying density field."
    },
    "0910/0910.4623_arXiv.txt": {
        "abstract": "We present a semi-analytical method to investigate the systematic effects and statistical uncertainties of the calculated angular power spectrum when incomplete spherical maps are used. The computed power spectrum suffers in particular a loss of angular frequency resolution, which  can be written as $\\delta \\ell \\sim \\pi/\\gamma_{max}$, where $\\gamma_{max}$ is the effective maximum extent of the partial spherical maps. We propose a correction algorithm to reduce systematic effects on the estimated $C_\\ell$, as obtained from the partial map projection on the spherical harmonic $\\ylm{\\ell}{m}$ basis. We have derived near optimal bands and weighting functions in $\\ell$-space for power spectrum calculation using small maps, and a correction algorithm for partially masked spherical maps that contain information on the angular correlations on all scales. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.3235_arXiv.txt": {
        "abstract": "We study both the background evolution and cosmological perturbations of anisotropic inflationary models supported by coupled scalar and vector fields. The models we study preserve the $U(1)$ gauge symmetry associated with the vector field, and therefore do not possess instabilities associated with longitudinal modes (which instead plague some recently proposed models of vector inflation and curvaton). We first intoduce a model in which the background anisotropy slowly decreases during inflation; we then confirm the stability of the background solution by studying the quadratic action for all the perturbations of the model. We then compute the spectrum of the $h_{\\times}$ gravitational wave polarization. The spectrum we find breaks statistical isotropy at the largest scales and reduces to the standard nearly scale invariant form at small scales. We finally discuss the possible relevance of our results to the large scale CMB anomalies. ",
        "introduction": "The inflationary paradigm has become a widely accepted description of the early universe, which has been successful in solving the classical problems of FRW cosmology~\\cite{Guth,Linde}. Moreover, inflation provides a mechanism for generation of nearly scale invariant spectrum of perturbations that can seed the structure we observe today~\\cite{mfb}. Inflationary era is characterized by a quasi de-Sitter expansion with a nearly constant Hubble rate. During inflation, quantum fluctuations are generated and amplified to scales above the Hubble scale (which specifies the horizon of casual processes) where they remain frozen until they re-enter the horizon after inflation has ended. A key assumption about the quantum fluctuations during the quasi de-Sitter stage is that they are in an adiabatic vacuum stage in the deep sub-horizon (early time) regime, which results in a nearly scale invariant spectrum.  Results of the WMAP experiment~\\cite{WMAP} are in overall agreement with the predictions of inflation. However, certain features of the full sky CMB maps seem to be anomalous in the standard picture. These anomalies include the low power in the quadrupole moment~\\cite{lowl}, the alignment of the lowest multipole moments (known as the 'axis-of-evil')~\\cite{axis}, and an asymmetry in power between the northern and southern ecliptic hemispheres~\\cite{asym}. The statistical significance of these anomalies has been extensively discussed in the literature and despite numerous efforts, an explanation due to a systematic effect or a foreground signal affecting the analysis is not forthcoming. These anomalies suggest a violation of statistical isotropy of cosmological fluctuations, which is at odds with standard mechanisms of generation of perturbations in inflationary models. Therefore, the possibility of relating the anomalies with modifications in the standard inflationary picture has been considered by numerous authors recently.  In the standard picture, $a_{lm}$ coefficients of CMB temperature anisotropies satisfy $< a_{lm}\\, a_{lm}^* > = C_l\\, \\delta_{ll'}\\, \\delta_{mm'}$ (so, the temperature fluctuations are said to be statistically isotropic). It has been shown in~\\cite{gcp,gcp2} that an anisotropic expansion that took place at the onset of standard single field inflation, leads to a nonvanishing correlation of the $a_{lm}$ coefficients for different $l$-modes. This violation of statistical isotropy might then be related to the alignment of the lowest multipoles observed in the CMB data. Therefore, several authors considered inflationary models with an anisotropic stage. As the simplest possible modification, anisotropic initial conditions at the onset of single field inflation have been considered in references~\\cite{gcp,gcp2,uzan1,uzan2}. Isotropization of the universe takes place due to the fast expansion supported by the inflaton field. Fluctuations at scales comparable to the horizon scale during the anisotropic stage of inflation are sensitive to the background evolution and thus provide breaking of statistical isotropy at those scales. If inflation had a limited duration, this could leave an imprint on the largest observable scales. On the other hand, fluctuations that leave the horizon after isotropization re-enter earlier than modes that left the horizon during the anisotropic stage, producing the standard nearly scale invariant spectrum at small scales~\\footnote{It was also shown in~\\cite{gkp} that the anisotropic stage could lead to a detectable gravitational wave signal.}. Since in such models isotropy is reached very soon (within one Hubble time) the breaking of statistical isotropy can be visible if the total duration of inflation is tuned to only the minimal duration to solve the standard problems of FRW cosmology.  The fine-tuning can be relaxed if the anisotropic stage is prolonged, due to the existence of other sources, for example vector fields. Recently, anisotropic expansion driven by vector fields have been studied by many authors. The anisotropic expansion is achieved when the vector field acquires a nonzero vacuum expectation value (VEV) along a spatial direction. There is a range of models which differ by the way the VEV is obtained. The first anisotropic inflationary model was considered long ago in reference~\\cite{ford}, where the $U(1)$ symmetry of the vector field action is broken by a potential term, which needs to have a negative curvature to support an anisotropic expansion. It has been shown in reference~\\cite{hcp1} that this model is plagued by a ghost instability, and its quantum theory is inconsistent~\\footnote{One may hope that this model has a well behaved UV completion. However, inflationary predictions are sensitive to this high energy regime, where the completion is needed.}. In another model, the ACW model~\\cite{acw}, the VEV is obtained by a lagrange multiplier field which fixes the norm of the vector field and breaks the $U(1)$ gauge symmetry. The stability analysis of the ACW model by considering the most general perturbations of the background has been performed in references~\\cite{dgw} and~\\cite{hcp1,hcp2}. In the latter studies, it has been shown that the linearized perturbations diverge close to horizon crossing, indicating the instability of the background solution. Another model~\\cite{soda1} considered a nonminimal coupling of the vector field to the scalar curvature which breaks the conformal invariance of the vector field action. The coupling with curvature allows for a slow-roll phase during the prolonged anisotropic expansion, and the universe isotropizes due to the existence of a massive scalar field, which is responsible for the overall isotropic expansion of the universe. Models of inflation with nonmininal coupling to curvature have also been considered for standard isotropic inflation~\\cite{mukhvect}. Isotropy of space is realized if $3N$ mutually orthogonal vector fields have equal VEVs. Instead, for randomly oriented $N$ vector fields, one expects a statistically isotropic background with order $1/\\sqrt{N}$ anisotropy. Perturbative calculations based around such a background configuration have been considered in references~\\cite{Golovnev:2008hv, Golovnev:2009ks, Golovnev:2009rm} and in~\\cite{hcp3}. The latter study is the only complete study linearized study of perturbations, taking into account all the physical degrees of freedom, and coupling between vector field and metric perturbations. It has been shown in~\\cite{hcp3} that the equations of motion for the linearized perturbations become singular close to horizon crossing, leading to instabilities. In all of these models (including~\\cite{acw,soda1,mukhvect}), instabilities are related to the existence of the longitudinal polarization of the vector field (which would otherwise be absent when the $U(1)$ symmetry is restored), which becomes a ghost close to horizon crossing (In the vector inflation model~\\cite{mukhvect}, ghosts also appear in the deep UV regime).  More recently, reference~\\cite{Watanabe:2009ct} introduced an anisotropic inflation model driven simultaneously by a vector field and a massive scalar field. The vector field is massless, but it is coupled to the scalar field through its kinetic term. Such type of coupling preserves the $U(1)$ gauge invariance of the vector field, so that the dangerous longitudinal vector polarization is absent; moreover the conformal invariance of the vector field is broken. When the vector field has a nonvanishing VEV along a spatial direction, the model possess an anisotropically expanding attractor solution. The anisotropy in the attractor is initially small, but it increases towards the end of inflation; therefore this is a counter example to the cosmic no hair conjecture (See~\\cite{Kaloper:1991rw} for a different example) .The breaking of conformal invariance due to the scalar-vector coupling can also be used to generate magnetic fields from inflation, as discussed in reference~\\cite{Turner:1987bw} and more recently in~\\cite{Gordon:2000hv, Demozzi:2009fu} and~\\cite{Watanabe:2009ct}. These models are expected to be free from ghost instabilities as long as the coupling to the scalar field remains positive. However to our knowledge, there has been no complete study of the stability of these models, which take into account all the physical degrees of freedom of the system (including gravity). A complete study of stability is therefore necessary, given the problems identified with other vector field models~\\cite{hcp1,hcp2,hcp3}.  In this work, we will present a complete study of cosmological perturbations of models with scalar-vector coupling that lead to an anisotropically expanding Bianchi-I background solution. We will also consider generalizations of the original model introduced in~\\cite{Watanabe:2009ct}, to include additional scalar fields, so that isotropization can be achieved before the end of inflation. This is required in order to obtain a scale invariant spectrum of perturbations at small scales, but the large scale spectrum is modified, which in turn can be related to the CMB anomalies. Our study has two main steps: one is to show that the model is free of ghost instabilities, and the second is to study the resulting phenomenology. We perform the first step explicitly in this paper. We develop the necessary tools to study the phenomenology of the model and as a simple exercise, we study only the spectrum of $h_{\\times}$ gravitational wave polarization in this paper.  The study of the full phenomenology (including the spectrum of the curvature perturbation) will be communicated elsewhere~\\cite{me}.  We will follow the formalism developed in~\\cite{gcp2,hcp2} to decompose and classify the perturbations. The generic background metric we study is the Bianchi-I metric with a residual two dimensional isotropy, given by \\begin{equation} ds^2 = -dt^2 + a(t)^2\\, dx^2 + b(t)^2\\, \\left( dy^2 + dz^2 \\right) \\nonumber \\end{equation} We exploit the two dimensional isotropy of the $y-z$ plane to decompose the perturbations into two decoupled classes: $2d$ scalar modes and $2d$ vector modes (As discussed in~\\cite{gcp2}, there is no two dimensional transverse-traceless mode). The linearized Einstein equations and the quadratic action for the two types of modes are decoupled and therefore we study them separately. We can deduce the number of degrees of freedom coming from each of the fields (gravity+vector field+scalar field) by a simple counting, which does not depend on the decomposition chosen to classify them. The metric has $10$ perturbations ($\\delta g_{\\mu\\nu}$) to start with, 4 of which can be removed by coordinate transformations. Of the remaining 6 modes, 4 are nondynamical, which can be best understood from the ADM formalism~\\cite{adm}. In the ADM formalism, the gravitational action is decomposed into the dynamical part containing the spatial metric $h_{ij}$ and a part containing the lapse (N) and shift functions ($N_i$). The lapse and shift functions have no kinetic terms in the action, and they can be integrated out by solving their equations of motion~\\footnote{The equations of motion derived from extremezing the gravitational action with respect to lapse and shift functions result in the momentum and hamiltonian constraints.}. This leaves only $h_{ij}$ as dynamical modes, which have only two degrees of freedom. In summary, the metric perturbations have only 2 dynamical degrees of freedom (which are the gravitational wave polarizations in the standard case) and 4 nondyncamical modes. For the case of the vector field, out of 4 perturbations to start with ($\\delta A_{\\mu}$), one perturbation can be removed by the $U(1)$ gauge transformation. Out of the remaining 3 perturbations 2 of them are dynamical and one mode is nondyamical ($\\delta A_0$). The scalar field perturbation introduces a single dynamical degree of freedom. Therefore, in total, the perturbations comprise of 5 dynamical modes, 2 of which are $2d$ vector and 3 are $2d$ scalar modes. When additional scalar fields are considered, each field introduces an extra scalar dynamical degree of freedom.  We will insert the perturbative expansion of the metric, the vector field and the scalar field(s) into the starting action and expand it at the quadratic order. We will show that the actions for the $2d$ vector and $2d$ scalar modes are decoupled, so we study them separately. We will determine the linear combinations of perturbations that canonically normalize the action, generalizing the computation of the standard isotropic case~\\cite{musa}. The study of the quadratic action is crucial for both showing that the model is consistent (free of ghosts) and also for determining the initial conditions for the perturbations. For modes that are inside the horizon, canonical combinations of perturbations can be quantized, and initial conditions can be set by the canonical commutation relations and by the requirement that the adiabatic vacuum state has minimal energy. We then proceed by numerically integrating the evolution equations for the canonical fields, starting from adiabatic initial conditions deep inside the horizon, until the end of inflation, which gives us the primordial power spectra. We do so for the $2d$ vector modes in this paper (which will be related to the power spectrum of the $h_{\\times}$ polarization of gravitational waves) and perform the study of the spectrum of $2d$ scalar modes in a separate publication.  The paper is organized as follows: In section~\\ref{sec:back}, we summarize the anisotropic inflationary background solution obtained in reference~\\cite{Watanabe:2009ct} and generalize the model to possess extra scalar fields. This generalization leads to isotropization before the end of inflation. In section~\\ref{sec:perturbations}, we discuss the perturbations around the background configuration. Specifically, in subsection~\\ref{sub:decomposition}, we review and discuss the decomposition of perturbations around the Bianchi-I background solution, following the formulation of reference~\\cite{gcp2}. In subsection~\\ref{sub:generic}, we discuss the generic properties of coupled vector-scalar models and develop tools that we will use in the following sections in order to find adiabatic initial conditions for such systems.  In subsection~\\ref{sub:2dV} we discuss $2d$ vector perturbations around the background configuration. We compute the action for $2d$ vector modes and find the combinations of perturbations that canonically normalize the action. In section~\\ref{sec:power}, we compute the power spectra of the $2d$ vector modes numerically. This study results in the spectrum of $h_{\\times}$ gravitational wave mode. We show that the spectrum has angular dependence (so breaks statistical isotropy) at large scales and reduces to the standard nearly scale invariant form at small scales. We also provide a fit to the numerically obtained spectrum. Finally in section~\\ref{sec:conclusions} we provide a general discussion about the results we have obtained and their possible relation to observations.  We also provide two extensive appendices: In appendix~\\ref{sub:app1}, we derive the relation between the $2d$ decomposed modes and the standard longitudinal mode, which has been used to determine the power spectra. In appendix~\\ref{sub:app2}, we perform the study of the $2d$ scalar modes. More precisely, we compute the quadratic action and determine the modes that canonically normalizes the action. The spectrum of scalar perturbations and the consequent phenomenological predictions (specifically, the $<a_{lm}\\, a_{l'm'}^*>$ correlation) will be presented elsewhere. ",
        "conclusions": "\\label{sec:conclusions}  In this work we have considered anisotropic inflationary models with coupled vector and scalar fields. The coupling of the scalar field to the vector field is chosen to preserve the $U(1)$ gauge symmetry of the Lagrangian, therefore, the longitudinal polarization of the vector field, which has been shown to cause instabilities~\\cite{hcp1,hcp2,hcp3}, does not exist. The anisotropic expansion is achieved when the vector field has a VEV along one of the spatial directions. The anisotropic solution is an attractor, with anisotropy increasing towards the end of inflation, so these types of models are counter examples to the cosmic no hair conjecture~\\cite{Wald}. We have considered a generalization of the original model introduced in reference~\\cite{Watanabe:2009ct}, to include additional scalar fields, and specifically concentrated on the double scalar field case. This modification was done in order to achieve isotropization of the universe before the end of inflation. The original model with a single scalar field coupled to the vector field has an anisotropic inflationary background which persists until the end of inflation. In this type of background, perturbations at all scales are modified. In the two field modification, the extra field is not coupled to the vector field, and it is responsible for the overall isotropic expansion of the universe only. We have chosen the ratio of the masses of the scalar fields such that the inflationary expansion takes place in two steps, the first being anisotropic and the second is isotropic. Therefore, only the largest scale perturbations are modified and at small scales the standard spectrum is recovered. As we have discussed in the introduction, vector fields with a nonvanishing VEV have been introduced since isotropization takes place very quickly in a Bianchi-I background, (when anisotropy is only coming from initial conditions) and thus leads to a fine tuning. We have introduced a two phase inflation, with the second phase being isotropic. The transition to the second isotropic phase is indeed very quick, however the fine tuning has been relaxed compared to the previous case. The second isotropic phase should still be tuned to last around $60$ e-folds, but this tuning is not so strict, since it only affects the scale where the power spectrum becomes statistically isotropic. Previously, the tuning was more stringent, since inflation must last for $60$ e-folds in order that the largest scales are modified from the initial anisotropic conditions.  A more important drawback of the inlfationary setup with only initial anisotropic conditions is that the initial singularity is very close to the time of isotropization, which is around one Hubble time. Thus, a calculation of the power spectrum shows that it becomes divergent at the largest scales~\\cite{gcp2}~\\footnote{This does not mean that the model is unstable, it reflects that nonlinear effects become important, and one has to go beyond linear perturbation theory.}.  We have studied the perturbations around the background configuration in detail, taking into account all the degrees of freedom of the system. Following the formalism developed in references~\\cite{gcp,gcp2}, we have classified decoupled sets of perturbations around the anisotropic background and studied them separately. Furthermore, we have computed the quadratic action for perturbations and showed that the model is stable and consistent. We also found linear combinations of the perturbations which canonically normalize the action. This enabled us to quantize the system which in turn determined the initial conditions for the perturbations in the adiabatic vacuum. We then numerically integrated the equations of motion for the $2d$ vector perturbations for a range of comoving momenta, which resulted in the power spectrum of the gravitational wave polarization $h_{\\times}$. We have only computed the quadratic action and determined the canonical modes for the $2d$ scalar perturbations. We leave the study of the power spectra and the consequent phenomenology, which in turn is related to the curvature and $h_+$ gravitational wave spectrum, to a later publication.  The computed spectrum for $h_{\\times}$ has certain interesting features, which might have some relevance to the observed CMB anomalies. First of all, the spectrum has angular dependence for $k<k_{iso}$ which breaks the statistical isotropy at large scales, therefore leads to an alignment of the lowest multipoles as described in~\\cite{gcp,gcp2,acw}. Moreover, the spectrum reduces to the standard nearly scale invariant form for $k>k_{iso}$, so the higher multipoles are not affected from the vector field. We have shown that the power spectrum is formally of the form $P_{\\times}(\\overrightarrow{k}) = P_{\\times,\\, iso}(k) + \\delta P_{\\times}(k,\\xi)$.  We expect that this will be true also for the spectrum of scalar modes. We have also provided a functional form for $\\delta P_{\\times}$ by fitting to the numerical spectrum in equation (\\ref{powerexp}). Phenomenologically acceptable modifications to the large scale spectrum should satisfy $\\vert \\delta P \\vert \\ll \\vert P_{iso} \\vert$, however, we have shown that for the model under discussion, this is not the case. Although our results are obtained for the gravitational wave mode $h_{\\times}$, we expect a similar behavior also for the curvature spectrum, which can be compared with observations. At larger scales, the power spectrum is greatly enhanced satisfying $\\vert \\delta P_{\\times} \\vert > \\vert P_{\\times, \\, iso} \\vert$. Although the anisotropy in the background is small (in the level of a percent, as can be seen from equation(\\ref{sigmadot})), the modification to the spectrum at large scales is not small. This is due to the fact that the background attractor solution is disconnected from the standard isotropic solution and therefore, the behavior of perturbations are dramatically different. There is no parameter in the attractor solution that connects it continuously to the isotropic FRW limit. The anisotropy is always small and proportional to $1/c^2$ and there is no finite limit in $c$ which can continuously deform the attractor solution to the isotropic solution. Therefore, the power spectrum cannot be approximated as to be the isotropic piece plus a small correction from anisotropy. This issue unfortunately limits the phenomenology of the model.  Even though the difficulties involved, anisotropic backgrounds still have potentially interesting phenomenological consequences. Off-diagonal correlations of the $a_{lm}$ coefficients of CMB temperature fluctuations arise in such backgrounds and can have relevance to the alignment of lowest multipoles. Moreover, at the largest scales, non standard scalar to tensor ratio is obtained, which can be tested by upcoming CMB experiments. The gravitational wave modes $h_{+}$ and $h_{\\times}$ behave differently at large scales, therefore a future detection of gravitational waves might be tested against the consequences of an early stage of anisotropic inflation. In this paper, we have provided a solid example of anisotropic inflationary calculation of perturbations and computed the spectrum of $h_{\\times}$ polarization of gravitational waves. Although the obtained spectrum has phenomenologically limited implications, we believe that our results are still relevant, since it is the first complete study of power spectrum in an anisotropic inflationary model driven by a vector field to our knowledge. In a further publication, we will study the spectrum of the curvature ${\\cal R}$ perturbation, which could then be compared with observational data."
    },
    "0910/0910.3690_arXiv.txt": {
        "abstract": "We describe the current status of solar modelling and focus on the problems originated with the introduction of solar abundance determinations with low CNO abundance values. We use models computed with solar abundance compilations obtained during the last decade, including the newest published abundances by Asplund and collaborators. Results presented here make focus both on helioseismic properties and the models as well as in the neutrino fluxes predictions. We also discuss changes in radiative opacities to restore agreement between helioseismology, solar models, and solar abundances and show the effect of such modifications on solar neutrino fluxes. ",
        "introduction": "}  Solar models are a corner stone of stellar astrophysics. The determination of the solar interior structure through helioseismology and, only a few years later, the discovery that neutrinos change flavor, gave spectacular confirmations of our ability to model the Sun and, by extension, of other stars. However, a series of works starting with a redetermination of the photospheric oxygen solar abundance \\citep{allende} and finishing with a complete revision of solar abundances \\citep{ags05}, led to a strong reduction in the overall metallicity of the Sun driven by much lower CNO and Ne abundances than previously determined. Soon after, solar models that adopted the new composition were shown to have an interior structure at odds with helioseismology determinations. Since then, the so-called {\\em solar abundance problem} has been in the spotlight of solar (and stellar) astrophysics. As nicely put by \\citet{dp06}, it represents the incompatibility between the best solar atmosphere and interior models available.  The effects of the low metallicity in the solar interior has been widely discussed in the literature. Among many other references, the reader can refer to \\citet{basu04,turck04,monta04,bs05,dp06} and \\citet{montecarlo}.  Some attempts to constrain the solar metallicity independently of photospheric measurements can be found in \\citet{antia06,lin07} and \\citet{bisonii}. The connection between solar neutrinos and composition has also been discussed in different works, e.g. \\citet{turck04,bp04,bs05,montecarlo,bps08,wick}. Possible modifications in the physical inputs of solar models have also been discussed in connection to the {\\em solar abundance problem}. The reader can refer to \\citet{monta04,guzik05,dp06,castro07} just to mention some relevant works.  In this article, we present a short review of the field and present new solar models that incorporate the most recent solar abundance determination by \\citet{agss09}. In \\S~\\ref{sec:models} we describe the main characteristics of the models used to obtain the core results presented here, including the different options for solar compositions we used. Results are presented in \\S~\\ref{sec:results} where helioseismology properties of the models and solar neutrino fluxes are discussed in the context of current observational and experimental data. In \\S~\\ref{sec:opac} we go to some length in discussing radiative opacities as a possible solution to the abundance problem, including the effects of opacities in solar neutrino fluxes. We summarize in \\S~\\ref{sec:conclu}. ",
        "conclusions": "}  We have attempted, in this incomplete review, to describe the current status of the {\\em solar abundance problem} that originated with new determinations of solar photospheric abundances from Asplund and collaborators. Results discussed here are based on models computed with both the original \\citep{ags05} and the newest \\citep{agss09} solar compositions. Our reference for a {\\em good} solar model is based on the \\citet{gs98} composition. The most important result is that with the new \\citep{agss09} abundances, the qualitative picture that emerged a few years ago, i.e. that low-Z solar models are in gross disagreement with helioseismology, remains the same. Quantitatively, the disagreement is less severe because the new abundances have slightly higher CNO abundances and a somewhat larger Ne abundance. The changes, however, do not help much neither in restoring the agreement with helioseismology nor in facilitating the way for alternative solutions in the form of modified input physics for solar models. We have described with some detail the effect of the new composition in opacities and the required change to recover good helioseismic properties. Changes of order 15\\% are needed, which are still much higher than currently estimated uncertainties in radiative opacities for the solar interior. In addition to helioseismic properties of the models, we have discussed the effects of the composition on the predicted neutrino fluxes and compared, when possible, with results from solar neutrino experiments. Additionally, we have tried to encourage efforts to experimentally determine neutrino fluxes from CNO bicycle, since these are the most sensitive fluxes to changes in abundances of CNO elements, thus offering the best chances for neutrinos to put direct constraints on the solar core composition."
    },
    "0910/0910.3373_arXiv.txt": {
        "abstract": "{ Multiple inflation is a model based on N=1 supergravity wherein there are sudden changes in the mass of the inflaton because it couples to `flat direction' scalar fields which undergo symmetry breaking phase transitions as the universe cools. The resulting brief violations of slow-roll evolution generate a non-gaussian signal which we find to be oscillatory and yielding $f_\\mathrm{NL} \\sim 5-20$. This is potentially detectable by e.g. Planck but would require new bispectrum estimators to do so. We also derive a model-independent result relating the period of oscillations of a phase transition during inflation to the period of oscillations in the primordial curvature perturbations generated by the inflaton.  } ",
        "introduction": "A major goal in modern observational cosmology is to detect any non-gaussianity of the temperature anisotropies in the cosmic microwave background (CMB) \\cite{Komatsu:2001rj,Bartolo:2004if,Chen:2006nt}. All single field, slow roll, inflation models give rise to anisotropies that are gaussian \\cite{Maldacena:2002vr}. Therefore detection of a significant non-gaussian signal would falsify such models and provide new insights into the dynamics of inflation.  Quantifying such a signal is however not a straightforward task. Whilst gaussianity is a well-defined property, non-gaussianity is not and the anisotropies can, in principle, deviate from gaussianity in many different ways \\cite{Chen:2006nt}. Also, the CMB temperature anisotropies measured by WMAP are \\emph{very} close to gaussian \\cite{Komatsu:2008hk}. Therefore any measure designed to detect non-gaussianity needs to be sensitive to a very small signal.  One measure of non-gaussianity that is particularly useful is the three point correlation function, or `bispectrum' of the temperature anisotropies. Gaussian random variables (and functionals of gaussian random variables) have the property that all odd power correlation functions are zero. This makes the bispectrum the lowest order statistic for which \\emph{any} non-zero result would indicate a departure from gaussianity. The bispectrum contains much more information than the power spectrum as, in general, it depends on both scale and shape. Therefore, if a non-zero bispectrum is detected it will also be an extremely useful statistic for constraining models of the early universe.  \\medskip However there are limitations to how much information can be inferred from the bispectrum:  \\begin{itemize}  \\item Modern CMB experiments possess very large numbers of pixels of data. e.g. $N \\simeq 10^6$ for WMAP and $N \\simeq 10^7$ for Planck \\cite{Ashdown:2006ey}. Considering that each higher order of correlation function will require an additional factor of $N$ calculations, exact determination of the bispectrum quickly becomes impossible. Therefore various estimators have been constructed \\cite{Fergusson:2008ra}, but each can look only for a specific type of bispectrum and so might miss a signal different from the type being searched for.  \\item Higher order correlation functions suffer more from cosmic variance because more information is required at each scale to compute them. In this work we will deal with a \\emph{scale dependent} signal, therefore limitations due to cosmic variance are particularly relevant.  \\end{itemize}  These problems are not necessarily insurmountable. Computing power will only continue to get stronger, making bispectrum estimators increasingly powerful with time. For each potential signal, estimators can also be constructed for that specific signal, ensuring processing time is used as efficiently as possible. The estimator can be optimised for detecting primordial non-gaussianity while discriminating against secondary non-gaussianities arising from e.g. unsubtracted point sources or residuals from component separation \\cite{Munshi:2009ik}. Finally, in models that have non-trivial scale dependence, a correlation will likely exist between the power spectrum and bispectrum as we illustrate in this paper.  One of the first methods for quantifying non-gaussianity involved rewriting the Newtonian potential $\\Phi(x)$ as \\cite{Komatsu:2001rj}, \\begin{equation} \\label{fNL local} \\Phi(x) = \\Phi_\\mathrm{L} (x) + f_\\mathrm{NL} (\\Phi_\\mathrm{L}^2 (x) - \\langle\\Phi_\\mathrm{L}^2 (x)\\rangle), \\end{equation} where $f_\\mathrm{NL}$ is a constant, and $\\Phi_\\mathrm{L} (x)$ is a gaussian variable. However, this parameterisation has a rather specific form and does not capture other possible deviations from gaussianity. A more general parameterisation can be obtained by allowing $f_\\mathrm{NL}$ to depend explicitly on the wavevector, ${\\bf k}$. In terms of the adiabatic curvature perturbation, $\\zeta$, the parameterisation becomes \\cite{Bartolo:2004if,Seery:2005wm}, \\begin{equation} \\label{fNL general} \\zeta(x) = \\zeta_\\mathrm{L} (x) - \\frac{3}{5} f_\\mathrm{NL}\\star(\\zeta_\\mathrm{L}^2 (x) - \\langle\\zeta_\\mathrm{L}^2 (x)\\rangle). \\end{equation} Here, the $\\star$ product is used because $f_\\mathrm{NL}$ has been allowed to depend on scale.\\footnote{See Eq.(229) in Ref.\\cite{Bartolo:2004if} for how to define this product.} The factor of -3/5 comes from the relationship between the adiabatic curvature and the Newtonian potential at matter domination.  For $f_\\mathrm{NL}$ as defined in Eq.(\\ref{fNL local}), the WMAP 5-year constraint is $-9 < f_\\mathrm{NL} < 111$ \\cite{Komatsu:2008hk}. This reinforces the point made already: the primordial temperature anisotropies are \\emph{very} close to gaussian. Given that the amplitude of these anisotropies is $\\Delta T/T \\sim 10^{-5}$ the non-gaussian contribution to these anisotropies is at most $\\sim10^{-8}$ (or $\\sim 0.1\\%$ of the overall anisotropy). In general the expectation from inflation is that the anisotropies should be close to gaussian, therefore this result can be considered as a success of the inflationary paradigm.  More precisely the expectation from inflation of non-gaussianity as parameterised by Eq.(\\ref{fNL general}) depends on the specific model. It is known that for single field, slow-roll inflation with a canonical kinetic term and a vacuum initial state, the result is $f_\\mathrm{NL} \\sim \\epsilon \\ll 1$ \\cite{Maldacena:2002vr} where $\\epsilon$ is the usual slow-roll parameter defined later in Eq.(\\ref{defeps}). The best sensitivity expected from the Planck satellite is to $f_\\mathrm{NL}$ of $\\mathcal{O}(5)$ while the secondary contribution from post-inflationary evolution is of $\\mathcal{O}(1)$ \\cite{Komatsu:2009kd}. Therefore the prediction of the simplest inflationary toy models~\\footnote{By this we mean a generic fine-tuned potential such as the frequently used $V (\\phi) = m^2 \\phi^2$ which has not yet been convincingly obtained from a physical theory, especially since inflation occurs in such models at $\\phi > M_\\mathrm{P}$.} is that there should \\emph{not} be a detection of primordial non-gaussianity in the near future.  Conversely a deviation from the simplest toy models can produce a larger value for $f_\\mathrm{NL}$. The study of these effects serves two purposes. Firstly, from the cosmological perspective, a detection of $f_\\mathrm{NL}$ will immediately rule out single field slow-roll inflation and focus attention on determining how inflation actually occured. Secondly, from the perspective of inflationary model building based on fundamental physics, as the bounds on $f_\\mathrm{NL}$ grow tighter, some interesting models can be ruled out if there is no detection.  Much work has been done on non-gaussianity generated due to multiple scalar fields (e.g. Ref.\\cite{Sasaki:2006kq}), non-canonical kinetic terms (e.g. Ref.\\cite{Langlois:2008qf}) or non-vacuum initial states (e.g. Ref.\\cite{Martin:1999fa}). However, surprisingly little attention has been paid to the possibility of non-gaussianity generated by a violation of slow-roll. Ref.\\cite{Byrnes:2009qy} did consider this in a context when multiple fields are present; however, the non-gaussianity itself is generated by the multiple fields, not the violation of slow-roll (see also Ref.\\cite{Cai:2009hw}). Refs.\\cite{Chen:2006xjb,Chen:2008wn} do consider the non-gaussianity generated by violation of slow-roll, but for a toy inflationary model. We follow their methods closely but consider a \\emph{physical} model of inflation. In both models the violation of slow-roll generates sharp features in the power spectrum and also generates a ringing in the bispectrum with a characteristic period which we calculate below.  \\subsection{Features in the Power Spectrum of Primordial Fluctuations}  It is an expectation in the simplest toy models of inflation that the power spectrum of temperature anisotropies in the CMB should be nearly scale-invariant \\cite{Lyth:1998xn}. However this need not be the case for physical models. If the observed spectrum was indeed perfectly scale-invariant, such models would be ruled out and na\\\"ively this might seem to be the case. Assuming the ``concordance'' $\\Lambda$CDM cosmology, the primordial spectrum is well parameterised as a power law with spectral index $n_\\mathrm{s} = 0.960 \\pm 0.014$ \\cite{Komatsu:2008hk}. This indicates that there cannot be any significant scale dependence that affects the spectrum over a large range of scales but it does \\emph{not} preclude the existence of e.g. localised oscillations in the spectrum, or other sharp features. Moreover since the observed anistrotopies arise from the convolution of the {\\em unknown} primordial spectrum with the transfer function of the {\\em assumed} cosmological model whose parameters are being determined, it is obvious that both unknowns cannot be determined simultaneously without further assumptions.  In fact, there are indications that the primordial spectrum might not be a scale-free power law, even assuming the $\\Lambda$CDM cosmology \\cite{Martin:2003sg,Kogo:2004vt,Shafieloo:2003gf,Tocchini-Valentini:2004ht,Nicholson:2009pi,Ichiki:2009zz}. At large angular scales ($>50^0$) there is essentially no power and there are anomalous `glitches', especially in the range of multipoles $\\ell \\simeq 20-40$ \\cite{Spergel:2003cb,Hinshaw:2006ia}. The statistical significance of these anomalies is not sufficient to claim a definite detection, however it is strong enough to provide some tension with the fit to a scale-free power-law primordial spectrum which has only a 3\\% probability of being a good description of the WMAP 1-year data \\cite{Spergel:2003cb}, although this did improve to $\\sim7\\%$ with the WMAP 3-year data release \\cite{Hinshaw:2006ia}. Future measurements of the $EE$ and $TE$ mode polarisations by the Planck and (proposed) CMBPol satellites will throw light on whether these glitches are real or not \\cite{Mortonson:2009qv}.  {\\it ``In the absence of an established theoretical framework in which to interpret these glitches (beyond the Gaussian, random phase paradigm), they will likely remain curiosities''} \\cite{Hinshaw:2006ia}. Indeed if there were a model that had predicted, without ambiguity, the position and amplitude of the glitches, this would be seen as very strong evidence for the model. Although no such model exists, the general possibility of generating glitches over a range of scales had in fact been proposed prior to the WMAP observations in the context of `multiple inflation' wherein the mass of the inflaton field undergoes sudden changes during inflation \\cite{Adams:1997de}. This generates characteristic localized oscillations in the spectrum, as was demonstrated numerically in a toy model of a inflationary potential with a `kink' parameterised as \\cite{Adams:2001vc}: \\begin{equation} V (\\phi) = \\frac{1}{2} m^2 \\phi^2 \\left[1 + c \\tanh \\left(\\frac{\\phi - \\phi_\\mathrm{s}}{d}\\right)\\right]. \\label{kinkpot} \\end{equation} By tuning the position, amplitude and gradient of the kink, the locations of the glitches can be varied to match the glitches seen in the power spectrum. One can then perform statistical likelihood tests to determine whether the fit is better with the glitches, but with the additional parameters, or with the simple scale-free spectrum. It was found by the WMAP team that the fit to the 1-year data improves significantly (by $\\Delta\\chi^2 = 10$) for the model parameters $\\phi_\\mathrm{s} = 15.5\\,M_\\mathrm{P}$, $c = 9.1\\times10^{-4}$ and $d = 1.4\\times10^{-2}\\,M_\\mathrm{P}$, where $M_\\mathrm{P}\\equiv(8\\pi\\,G_\\mathrm{N})^{-1/2} \\simeq 2.44\\times10^{18}$~GeV \\cite{Peiris:2003ff}. This analysis was repeated later using the WMAP 3-year data, with similar results \\cite{Covi:2006ci}.  This seems encouraging, however $m$ in the toy model above is {\\em not} the mass of the inflaton --- in fact in all such monomial `chaotic' inflation models with $V \\propto \\phi^n$, inflation occurs at field values $\\phi_\\mathrm{infl} > M_\\mathrm{P}$, hence the leading term in a Taylor expansion of the potential around $\\phi_\\mathrm{infl}$ is always linear in $\\phi$ (since this is not a point of symmetry), rather than quadratic as for a mass term \\cite{German:1999gi}. The effect of a change in the inflaton mass can be sensibly modelled only in `new' inflation where inflation occurs at field values $\\phi_\\mathrm{infl} << M_{\\rm Pl}$ and an effective field theory description of the inflaton potential is possible. The `slow-roll' conditions are violated when the inflaton mass changes due to its (gravitational) coupling to `flat direction' fields which undergo thermal phase transitions as the universe cools during inflation \\cite{Adams:1997de}. The resulting effect on the spectrum of the curvature perturbation was found by analytic solution of the governing equations to correspond to a `step' followed by rapidly damped oscillations \\cite{Hunt:2004vt}.\\footnote{Although a similar phenomenon had been noted earlier for the case where the inflaton potential has a jump in its slope \\cite{Starobinsky:ts}, such a discontinuity has no physical interpretation.}  The next step should be to predict other observables, having used the power spectrum to constrain all the parameters in the model. In this paper we calculate the bispectrum of the multiple inflation model \\cite{Adams:1997de}, using its predicted power spectrum \\cite{Hunt:2004vt} and the set of parameters which provide the best reduced $\\chi^2$ in the $\\Lambda$CDM cosmology \\cite{Hunt:2007dn}. We also examine the effect on the bispectrum of varying these parameters over their full natural range.\\footnote{We use the word `natural' in this context to mean ``stable towards radiative corrections''.}  The non-gaussianity is found to be potentially detectable by the Planck satellite, or perhaps even a reanalysis of the WMAP data. ",
        "conclusions": "There is tentative evidence that the primordial power spectrum of scalar perturbations is imprinted with sharp features on large scales \\cite{Spergel:2003cb,Hinshaw:2006ia}. These could have resulted from one or more phase transitions early on in the inflationary epoch as happens in the multiple inflation model \\cite{Adams:1997de}. This model is well motivated by fundamental physics ($N=1$ supergravity) and such spectral features were predicted \\emph{before} WMAP provided observational indications for them.  It was shown \\cite{Chen:2006xjb} that any such departure from slow-roll during inflation should also generate non-gaussianity. We have calculated the non-gaussianity in mutliple inflation for the parameter values that best fit the features in the power spectrum \\cite{Hunt:2007dn}. This is on the edge of being observable but a clear detection would require a new type of bispectrum estimator due to its scale dependence and non factorisability \\cite{Fergusson:2008ra}. Significantly larger non-gaussianities can be generated during multiple inflation if the parameters are allowed to range over values which are technically natural (i.e. stable towards radiative corrections).  The form of this non-gaussianity is tightly correlated with the power spectrum --- the oscillations in the bispectrum should begin at the same multipole as the oscillations in the power spectrum and have two thirds of the period. There have been a number of attempts, to deconvolve the primordial power spectrum directly from the WMAP data \\cite{Kogo:2004vt,Shafieloo:2003gf,Tocchini-Valentini:2004ht} (see also Ref.\\cite{Bridges:2005br}). It is generally found that there is a supression of power on the scale of the present Hubble radius, followed by a `ringing' at medium scales. By measuring the period of the oscillations in the power spectrum one can predict the period (and phase) of oscillations in the bispectrum, in a \\emph{model independent} manner.  Forthcoming measurements of CMB polarisation by Planck ought to shed light on whether these features in the power spectrum are systematic errors or genuine evidence of non-trivial dynamics during the inflationary era. The associated non-gaussian signal should also be detectable according to our model calculation and provide insight into the dynamics.  \\subsection{Acknowledgements}  SH acknowledges a Clarendon Fellowship and support from Balliol College, Oxford and the EU Marie Curie Network ``UniverseNet'' (HPRN-CT-2006-035863). We thank Paul Hunt, David Lyth, Graham Ross, Misao Sasaki and David Wands for helpful comments and discussions."
    },
    "0910/0910.4765_arXiv.txt": {
        "abstract": "The next decade will bring massive new data sets from experiments of the direct detection of weakly interacting massive particle (WIMP) dark matter.  Mapping the data sets to the particle-physics properties of dark matter is complicated not only by the considerable uncertainties in the dark-matter model, but by its poorly constrained local distribution function (the ``astrophysics'' of dark matter).  I propose a shift in how to think about direct-detection data analysis.  I show that by treating the astrophysical and particle-physics uncertainties of dark matter on equal footing, and by incorporating a combination of data sets into the analysis, one may recover both the particle physics and astrophysics of dark matter.  Not only does such an approach yield more accurate estimates of dark-matter properties, but it may illuminate how dark matter coevolves with galaxies. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.1971_arXiv.txt": {
        "abstract": "The \\amidas\\ website has been established as an online interactive tool for running simulations and analyzing data in direct Dark Matter detection experiments. At the first phase of the website building, only some commonly used WIMP velocity distribution functions and elastic nuclear form factors have been involved in the \\amidas\\ code. In order to let the options for velocity distribution as well as for nuclear form factors be more flexible, we have extended the \\amidas\\ code to be able to include {\\em user--uploaded} files with their own functions. In this article, I describe the preparation of files of user--defined functions onto the \\amidas\\ website. Some examples will also be given. ",
        "introduction": "In the last few years we developed new methods for analyzing data, i.e., measured recoil energies, from (future) direct Dark Matter detection experiments as model--independently as possible \\cite{DMDDf1v, DMDDmchi, DMDDfp2-IDM2008, DMDDidentification-DARK2009}. These methods will help us to understand the nature of WIMP (Weakly Interacting Massive Particle $\\chi$) Dark Matter, to identify them among new particles produced hopefully in the near future at colliders, as well as to reconstruct the (sub)structure of our Galactic halo. Following the development of these model--independent data analysis procedures, we combined the programs for simulations to a compact system: \\amidas\\ (A Model--Independent Data Analysis System). For users' convenience and under the collaboration with the ILIAS Project \\cite{ILIAS}, an online system has also been established at the same time \\cite{AMIDAS-web, AMIDAS-SUSY09}.  For the first version of the \\amidas\\ code and website, the options for target nuclei, for the velocity distribution function of halo WIMPs, as well as for the elastic nuclear form factors for spin--independent (SI) and spin--dependent (SD) WIMP--nucleus interactions are fixed and only some commonly used forms have been involved in the \\amidas\\ code \\cite{AMIDAS-SUSY09}. Users can not choose different detector materials nor use different velocity distribution/form factors for their simulations and/or data analyses. In order to let the options for the velocity distribution as well as for the nuclear form factors be more flexible, we have extended the \\amidas\\ code to be able to include {\\em user--uploaded} files with their own functions. Note that, since the \\amidas\\ code has been written in the C programming language, all user--defined functions for uploading must also be given using the syntax of C.  The remainder of this article is organized as follows. In Sec.~2 I will talk about setting users' own target nuclei. In Secs.~3 and 4 the preparation of files defining the velocity distribution function and the nuclear form factors will be described, respectively. I will conclude in Sec.~5. Some intrinsically defined constants and functions in the \\amidas\\ code will be given in an appendix. ",
        "conclusions": "In this article, I described the preparation of files giving user--defined WIMP velocity distribution function and/or elastic nuclear form factor(s) which can be uploaded onto the \\amidas\\ website for more flexible simulations or data analyses. This improvement allows theorists to simulate with their own/favorite models and compare them with (future) experimental results, as well as gives experimentalists flexible choices for more suitable form factor(s) for their own detector materials.  In summary, up to now all basic functions of the \\amidas\\ code and website have been well established. Hopefully this new tool can help our colleagues to detect/discover WIMP Dark Matter, to understand the nature of Dark Matter particles and the (sub)structure of the Galactic halo in the future. \\subsubsection*"
    },
    "0910/0910.4881_arXiv.txt": {
        "abstract": "We report on observations of TeV-selected AGN made during the first 5.5 months of observations with the Large Area Telescope (LAT) on-board the \\textit{Fermi Gamma-ray Space Telescope} (\\Fermic). In total, \\NObjTotal\\ AGN were selected for study, each being either (i) a source detected at TeV energies (\\NObjTeVSrc\\ sources) or (ii) an object that has been studied with TeV instruments and for which an upper-limit has been reported (\\NObjTeVLim\\ objects). The \\Fermi observations show clear detections of \\NDetTotal\\ of these TeV-selected objects, of which \\NDetTeVSrc\\ are joint GeV--TeV sources and \\NDetTotalNoEGRET\\ were not in the third EGRET catalog. For each of the \\NDetTotal\\ \\Fermi-detected sources, spectra and light curves are presented. Most can be described with a power law of spectral index harder than 2.0, with a spectral break generally required to accommodate the TeV measurements. Based on an extrapolation of the \\Fermi spectrum, we identify sources, not previously detected at TeV energies, which are promising targets for TeV instruments. Evidence for systematic evolution of the \\gray spectrum with redshift is presented and discussed in the context of interaction with the EBL. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.2375_arXiv.txt": {
        "abstract": "Computer representations of real numbers are necessarily discrete, with some finite resolution, discreteness, quantization, or minimum representable difference.  We perform astrometric and photometric measurements on stars and co-add multiple observations of faint sources to demonstrate that essentially all of the scientific information in an optical astronomical image can be preserved or transmitted when the minimum representable difference is a factor of two finer than the root-variance of the per-pixel noise. Adopting a representation this coarse reduces bandwidth for data acquisition, transmission, or storage, or permits better use of the system dynamic range, without sacrificing any information for down-stream data analysis, including information on sources fainter than the minimum representable difference itself. ",
        "introduction": "Computers operate on bits and collections of bits; the numbers stored by a computer are necessarily discrete; finite in both range and resolution.  Computer-mediated measurements or quantitative observations of the world are therefore only approximately real-valued.  This means that choices must be made, in the design of a computer instrument or a computational representation of data, about the range and resolution of represented numbers.  In astronomy this limitation is keenly felt at the present day in optical imaging systems, where the analog-to-digital conversion of CCD or equivalent detector read-out happens in real time and is severely limited in bandwidth; often there are only eight bits per readout pixel.  This is even more constrained in space missions, where it is not just the bandwidth of real-time electronics but the bandwidth of telemetry of data from space to ground that is limited.  If the ``gain'' of the system is set too far in one direction, too much of the dynamic range is spent on noise, and bright sources saturate the representation too frequently.  If the gain is set too far in the other direction, information is lost about faint sources.  Fortunately, the information content of any astronomical image is limited \\emph{naturally} by the fact that the image contains \\emph{noise}.  That is, tiny differences between pixel values---differences much smaller than the amplitude of any additive noise---do not carry very much astronomical information.  For this reason, the discreteness of computer representations of pixel values do not have to limit the scientific information content in a computer-recorded image.  All that is required is that the noise in the image be \\emph{resolved} by the representation.  What this means, quantitatively, for the design of imaging systems is the subject of this \\documentname; we are asking this question: ``What bandwidth is required to deliver the scientific information content of a computer-recorded image?''  This question has been asked before, using information theory, in the context of telemetry \\citep{Gaztanaga} or image compression \\citep{Watson}, treating the pixels (or linear combinations of them) as independent.  Here we ask this question, in some sense, \\emph{experimentally}, and for the properties of imaging on which optical astronomy depends, where groups of contiguous pixels are used in concert to detect and centroid faint sources.  We perform experiments with artificial data, varying the bandwidth of the representation---the size of the smallest representable difference $\\Delta$ in pixel values---and measuring properties of scientific interest in the image.  We go beyond previous experiments of this kind (\\citealt{WhiteGreenfield}, and \\citealt{PenceInPress}) by measuring the centroids and brightnesses of compact sources, and sources fainter than the detection limit.  The higher the bandwidth, the better these measurements become, in precision and in accuracy. We find, in agreement with previous experiments and information-theory-based results, that the smallest representable difference $\\Delta$ should be on the order of the root-variance $\\sigma$ of the noise in the image.  More specifically, we find that \\emph{the minimum representable difference should be about half the per-pixel noise sigma} or that about two bits should span the FWHM of the noise distribution if the computer representation is to deliver the information content of the image.  Of course, tiny mean differences in pixel values, even differences much smaller than the noise amplitude, \\emph{do} contain \\emph{extremely valuable} information, as is clear when many short exposures (for example) of one patch of the sky are co-added or analyzed simultaneously.  ``Blank'' or noise-dominated parts of the individual images become signal-dominated in the co-added image.  In what follows, we explicitly include this ``below-the-noise'' information as part of the information content of the image.  Perhaps surprisingly, \\emph{all} of the information can be preserved, even about sources fainter than the discreteness of the computer representation, provided that the discreteness is finer than the amplitude of the noise.  This result has important implications for image compression, but our main interest here is in the design and configuration of systems that efficiently take or store raw data, using as much of the necessarily limited dynamic range on signal as possible.  Our results have some relationship to the study of \\emph{stochastic resonance}, where it has been shown that signals of low dynamic range can be better detected in the presence of noise than in the absence of noise (see \\citealt{stochres} for a review).  These studies show that if a signal is below the minimum representable difference $\\Delta$, it is visible in the data only when the digitization of the signal is noisy.  A crude summary of this literature is that the optimal noise amplitude is comparable to the minimum representable difference $\\Delta$.  We turn the stochastic resonance problem on its head: The counterintuitive result (in the stochastic resonance context) that weak signals become detectable only when the digitization is noisy becomes the relatively obvious result (in our context) that so long as the minimum representable difference $\\Delta$ is comparable to or smaller than the noise, signals are transmitted at the maximum fidelity possible in the data set.  What we call here the ``minimum representable difference'' has also been called by other authors the ``discretization'' \\citep{Gaztanaga} or ``quantization'' \\citep{Watson, WhiteGreenfield, Pence}. ",
        "conclusions": "Because of finite noise, the information content in astronomical images is finite, and can be captured by a finite numerical resolution.  In the above, we scaled and snapped-to-integer real-valued images by a SNIP procedure such that in the SNIPped image, the minimum representable difference $\\Delta$ between pixel values was set to a definite fraction of the Gaussian noise root-variance (sigma) $\\sigma$.  We found with direct numerical experiments that the SNIP procedure introduces essentially no significant error in estimating the variance of the image, or in centroiding or photometering stars in the image, when the minimum representable difference is set to any value $\\Delta \\leq 0.5\\,\\sigma$.  In addition, we showed that all the information about sources fainter than the per-pixel noise level is preserved by the quantization (SNIP) procedure, again provided that $\\Delta \\leq 0.5\\,\\sigma$.  This is somewhat remarkable because at $\\Delta = 0.5\\,\\sigma$ the faintest sources in our experiments were fainter than the minimum representable difference.  Although it is somewhat counterintuitive that integer quantization of the data does not remove information about sources fainter than the quantization level, it is perhaps even more counterintuitive how well photometric measurements perform in our coadd tests.  For example, in \\figurename s~\\ref{fig:bitsoffsetcoadd1} and \\ref{fig:bitsoffsetcoadd2}, the photometric measurements are relatively accurate even when the data are quantized at minimum representable difference $\\Delta = 16\\,\\sigma$!  The quality of the measurements can be understood in part by noting that the coadded images have a per-pixel noise $\\sigma = \\sqrt{1024}\\,\\sigma = 32\\,\\sigma$, which is once again larger than the minimum representable difference, and in part by noting that each image has a different sky level, so each individual image is differently ``wrong'' in its photometry; many of these differences average out in the coadd.  When the sky level is held fixed across coadded images, photometric measurements become inaccurate again---as seen in \\figurename~\\ref{fig:bitsoffsetcoadd1_samesky}---because individual-image biases caused by the coarse quantization no longer ``average out''.  Our fundamental conclusion is that all of the scientifically relevant information in an astronomical image is preserved as long as the minimum representable difference $\\Delta$ in pixel values is smaller than or equal to half the per-pixel root-variance (sigma) $\\sigma$ in the image noise.  This confirms previous results based on information-theory arguments (for example, \\citealt{Gaztanaga}), and extends previous experiments on bright-source photometry \\citep{WhiteGreenfield, PenceInPress} to astrometry and to sources fainter than the noise.  Our experiments were performed on images with pure Gaussian noise; of course many images contain significant non-Gaussianity in their per-pixel noise so the empirical variance will depart significantly from the true noise variance \\citep{WhiteGreenfield}.  The conservative approach for such images is to take not the true variance for $\\sigma^2$ but rather use for $\\sigma^2$ something like the minimum of the straightforwardly measured variance and a central variance estimate, such as a sigma-clipped variance estimate, an estimate based on the curvature of the central part of the noise value frequency distribution function, or the median absolute difference of nearby pixels \\citep{Pence}.  With this re-definition of the root variance $\\sigma$, the condition $\\Delta \\leq 0.5\\,\\sigma$ represents a conservative setting of the minimum representable difference.  The fact that a $\\Delta = 0.5\\,\\sigma$ representation preserves information on the faint sources---even those fainter than $\\Delta$ itself---has implications for the design of data-taking systems, which are necessarily limited in bandwidth.  If the system is set with $\\Delta$ substantially smaller than $0.5\\,\\sigma$, then bright sources will saturate the representation more frequently than necessary, while no additional information is being carried about the faintest sources. Any increase in $\\Delta$ pays off directly in putting more of the necessarily limited system dynamic range onto bright sources, so it behooves system designers to push as close to the $\\Delta = 0.5\\,\\sigma$ limit as possible.  To put this in the context of a real data system, we looked at a ``DARK'' calibration image from the Hubble Space Telescope Advanced Camera for Surveys (ACS).  The dark image should have the lowest per-pixel noise of any ACS image, because it has only dark and read noise.  We chose image set \\texttt{jbanbea2q}, and measured the median noise level in the raw DARK image with the median absolute difference between values of nearby pixels (for robustness).  The ACS data system is operating with a minimum representable difference $0.25\\,\\sigma < \\Delta < 0.33\\,\\sigma$, comfortably within the information-preserving range and close to the minimum-bandwidth limit of $\\Delta = 0.5\\,\\sigma$.  Of course this is for a dark frame; sky exposures (especially long ones) could have been profitably taken with a larger $\\Delta$ (because $\\sigma$ will be greater); this would have preserved more of the system dynamic range for bright sources.  If the ACS took almost exclusively long exposures, the output would contain more scientific information with a larger setting of the minimum representable difference.  In some sense, the results of this paper recommend a ``lossy'' image compression technique, in which data are scaled by a factor and snapped to integer values such that the minimum representable difference $\\Delta$ is made equal to or smaller than $0.5\\,\\sigma$. Indeed, when typical real-valued astronomical images are converted to integers at this resolution, the integer versions compress far better with subsequent standard file compression techniques (such as gzip) than do the floating-point originals \\citep{Gaztanaga, Watson, WhiteGreenfield, Pence, bernstein}.  In the $\\Delta = 0.5\\,\\sigma$ representation, after lossless compression, storage and transmission of the image ``costs'' only a few bits per noise-dominated pixel. Because the snap-to-integer step changes the data, this overall procedure is technically lossy, but we have shown here that none of the \\emph{scientific} information in the image has been lost."
    },
    "0910/0910.5280_arXiv.txt": {
        "abstract": "This is my contribution to Proceedings of the International Workshop on Cosmic Structure and Evolution, September 23-25, 2009, Bielefeld , Germany. In my talk I presented some non-Gaussian features of the foreground reduced WMAP five year full sky temperature maps, which were recently reported in the Ref. \\cite{Vitaly}. And in these notes I first discuss the statistics behind this analysis in some detail. Then I describe invaluable insights which I got from discussions after my talk on the Workshop. And finally I explain why, in my current opinion, the signal detected in the Ref. \\cite{Vitaly} can hardly have something to do with cosmological perturbations, but rather it presents a fancy measurement of the Milky Way angular width in the microwave frequency range. ",
        "introduction": "Nowadays we witness a great progress in both theoretical and observational cosmology which makes our demands and expectations ever higher and turns us to discussing more and more subtle properties of the available data. One of such popular topics is the quest for primordial non-Gaussianities in the spectrum of the CMB radiation. Not really expected to be detectable for the simplest models of inflation, these small departures from the purely Gaussian signal would help to distinguish between more elaborate inflationary scenarios and would probably provide us with some new insights into the wonderful realm of the very early Universe.  The approach I discuss is based on a very simple idea. Assume that the Universe is statistically isotropic, and all the temperature fluctuations in the CMB radiation are of statistical nature. Then we decompose the fluctuations, as usual, into the spherical harmonics denoting the coefficients by $a_{l,m}$ and get $$\\langle a_{l,m}a_{l^{\\prime},m^{\\prime}}^{*}\\rangle=\\langle a_{l,m}a_{l^{\\prime},-m^{\\prime}}\\rangle =C_l\\delta_{l,l^{\\prime}}\\delta_{m,m^{\\prime}}.$$ Moreover, $a_{l,m}$'s with a fixed value of $l$ but different values of $m$ can be thought of as different realizations of one and the same random variable. So that one can naturally ask a question about the shape of the distribution. It can be answered by many methods, from plotting a histogram of the sample to estimating the higher moments of the distribution. A Gaussian distribution is completely determined by two parameters, its mean and its variance. For fluctuations the mean is taken to be $0$, and the only parameter left is the variance, $C_l$. (Of course, it implies rescaled $\\chi^2$-distributions for quadratic in $a$ quantities.) If the random variable is known (or assumed) to be Gaussian, this single parameter can, in principle, be extracted from any part of the probability distribution. The idea of the Ref. \\cite{Vitaly} is to take only the tails, i.e. to deduce the variance $C_l$ once more from only the distribution of large coefficients, $|a_{l,m}|^2>C_l$, and then to compare this result with the original one. Up to statistical variations in the number of data points in the tails, this corresponds to using the order statistics with somewhat more points than in top and bottom sextiles but with less points than in two marginal quintiles.  Applied to the full sky foreground reduced maps, this method gives a well-pronounced peak in the difference between the two estimates for the variance. The peak is located in the range of $l\\approx 45 \\pm 15$. The fluctuations outside of the peak are also much larger than would be expected statistically. This is, of course, due to remaining foregrounds contamination, and the only reason to take the peak seriously was that it is a few more times larger than the other fluctuations \\cite{Vitaly} and it looks more or less the same in different frequency bands (up to the different overall level of noise). And this was also my conclusion that it should have something to do with cosmology. I explain the relevant statistics in Sections 2 and 3. And for the graphical presentation of results, I refer the reader to \\cite{Vitaly}. However, at the Workshop I have learned from Pavel Naselsky that the multipole coefficients with even values of $l+m$ have the worst contamination from Galaxy, see below. In Section 4 I discuss the data analysis with separation of $(l+m)$-even and $(l+m)$-odd harmonics, and show that the initial assumption of having different realizations of one random variable is heavily disproved due to Galactic signal which invalidates the claim for cosmological non-Gaussianities. My current conclusion presented in the Section 5 is that the effects observed in \\cite{Vitaly} refer to the structure of Galactic noise, and not to properties of primordial fluctuations. Due to this reason I cancelled my authorship for the second version of that article. (I was a co-author for the first one). And I would like to note here that all the computer work for the article \\cite{Vitaly} was done by my former co-author Vitaly Vanchurin and, needless to say, if there appears to be something primordial about this peak then the whole success should be attributed solely to him and his enthusiasm. An interested reader may also want to consult with the original reference \\cite{Vitaly} for the opinion opposite to mine. ",
        "conclusions": "In these notes I discussed the statistics behind the non-Gaussian anomalies reported in Ref. \\cite{Vitaly}. After that I have shown that the most probable explanation of the signal refers to geometric properties of the Galactic noise, and not to cosmology. Admittedly, I do not have a good understanding of the structure of Galactic noise. But at the very least, the claim for cosmological non-Gaussianities is pretty much premature. (I refer an interested reader to the work \\cite{Vitaly} for a different opinion.) On the other hand, one could probably use this kind of analysis to extract some information about the structure of the noise.  {\\bf Acknowledgments.} I am very grateful to the organizers of the Workshop for the opportunity to participate in this wonderful event, and especially I wish to thank Dominik Schwarz for invitation and for his encouragement to write this contribution. I am also very grateful to Pavel Naselsky and to other participants for very useful discussions. This work was supported in part by the Cluster of Excellence EXC 153 {}``Origin and Structure of the Universe''."
    },
    "0910/0910.0831_arXiv.txt": {
        "abstract": "We report the first interferomteric detection of 183 GHz water emission in the low-mass protostar Serpens SMM1 using the Submillimeter Array with a resolution of 3$''$ and rms of $\\sim$7 Jy in a 3 km s$^{-1}$ bin. Due to the small size and high brightnessof more than 240 Jy/beam, it appears to be maser emission. In total three maser spots were detected out to $\\sim$ 700 AU from the central protostar, lying along the red-shifted outflow axis, outside the circumstellar disk but within the envelope region as evidenced by the continuum measurements. Two of the maser spots  appear to be blue-shifted by about 1 to 2 km s$^{-1}$. No extended or compact thermal emission from a passively heated protostellar envelope was detected with a limit of 7 Jy (16 K), in agreement with recent modelling efforts.  We propose that the maser spots originate within the cavity walls due to the interaction of the outflow jet with the surrounding protostellar envelope. Hydrodynamical models predict that such regions can be dense and warm enough to invert  the 183 GHz water transition. ",
        "introduction": "Water is one of the most important molecules in interstellar clouds in general and in star-forming regions in particular. In warm regions ($T>$100 K) close to the protostar, water is prominent with gas abundances up to 3$\\times10^{-4}$ w.r.t. \\HH, even higher than CO \\citep{Cernicharo90,vanDishoeck96,Harwit98}. Besides the inner regions of protostellar envelopes these high abundances are also found where powerful jets from the protostar interact with the surroundings \\citep{Nisini99}. In contrast, water abundances are as low as $10^{-8}$-$10^{-9}$ in cold ($T<$100 K) envelope regions as evidenced by Infrared Space Observatory (ISO) and Submillimeter Wave Astronomy Satellite (SWAS) observations \\citep[e.g][]{Boonman03}. Emission from  water molecules is very difficult to observe, due to the limits imposed by the Earth's atmosphere. Isotopologues such as deuterated water can be observed from the ground \\citep[e.g.][]{Schulz91,Parise05} The only transitions of the main water isotope that can be observed regularly are low-frequency maser transitions. Most famous is the water maser at 22.2 GHz, the 6$_{16}$-5$_{23}$ transition, regularly observed using radio telescopes and interferometers, such as the Very Large Array (VLA). In star forming regions with embedded sources of high luminosity ($>$100 L$_{\\odot}$) 22.2 GHz water maser emission is commonly detected and used to probe gas kinematics \\citep[e.g.][]{Moscadelli06,Goddi06}. The emission is found to be variable on timescales of a day to a month, probably related to variations in the accretion disk or the outflowing material \\citep[e.g.][]{Pashchenko05}. In regions with embedded sources of lower luminosity, 22.2 GHz water maser emission is detected less frequently than in high-mass sources, but again with considerable variability  \\citep[e.g.][]{Wilking94,Claussen96,Claussen98,Furuya03}. Imaging with the VLA detected the  maser emission within several hundred AU of low-mass protostars \\citep[e.g][]{Furuya99,Furuya01}, while Very Long Baseline Array observations indicate that the location may even be closer to the protostar \\citep[e.g][]{Moscadelli06}. Both water and methanol masers have been observed to be related to disks \\citep{Torrelles96,Moscadelli06}, while other observations of masers have been associated with outflows \\citep[e.g.][]{Claussen96,Furuya03}. However, the excitation conditions for the 22 GHz water maser line are very high density ($>$10$^8$ cm$^{-3}$) and temperature (2000 K$>$ T $>$200 K)\\citep{Yates97}, a combination of conditions that is rare in low-mass protostars.  Another water maser transition is found at 183.3 GHz. This 3$_{13}$-2$_{20}$ transition has been used in extra-galactic and galactic studies, primarily with the 30m telescope at Pico Veleta, Spain \\citep[e.g.][]{Cernicharo94,Cernicharo96}. Statistical equilibrium calculations combined with the observation of extended emission conclude that relatively low temperatures (T$\\sim$ 150 K) and densities (10$^5$-10$^6$ cm$^{-3}$) can already invert the populations of the 183.3 GHz transition \\citep{Cernicharo94}. In star forming regions it has been detected in several low- and high-mass protostellar sources such as Orion, W49N, H7-11 and L1448-mm. Most lines consist of a broad component,  superposed with a strong narrow line, presumably a maser line. The broad component was found to be spatially extended thermal emission after having been mapped in some high-mass star forming regions such as Orion and W49N \\citep{Cernicharo94,GonzalesAlfonso95}. Extended thermal emission was also found  in low-mass sources such as HH7-11 \\citep{Cernicharo96}.   Serpens SMM1 (referred to as SMM1) is a low-mass Class 0 source in the Serpens cluster ($D$= 250 pc) and has been studied at (sub)millimeter wavelengths by \\citet{Hogerheijde99}. This source was selected as part of a pilot study to observe the 183 GHz line due to its inclusion in the Water in Star-forming regions with Herschel (WISH) program\\footnote{See http://www.strw.leidenuniv.nl/WISH} as a source where many water lines will be targetted with spatial resolutions of 9-40$''$. It is relatively luminous ($L$=20.7 $L_\\odot$) and thus an ideal target for water observations. It was observed with ISO-LWS \\citep{Larsson02}, where numerous thermal water lines were found in the large 120$''$ beam. It drives a powerful highly collimated radio jet \\citep{Curiel93}, a large molecular outflow \\citep{White95} and was included in the 22 GHz water maser survey by \\citet{Furuya03}, where several maser spots were found. These masers were in turn studied by \\citet{Moscadelli06} who found the masers originating within a circumstellar disk based on their position and proper motions.  Observations with SWAS (beam = 3.3$'$ by 4.5$'$) of the 1$_{10}$-1$_{01}$ water line found water associated with the outflow with an abundance of o-H$_2$O of a few times 10$^{-7}$ \\citep{Franklin08} in both red and blue outflow lobes.  In this letter, we present first interferometric  observations using the SMA\\footnote{The Submillimeter Array is a joint project between the Smithsonian Astrophysical Observatory and the Academia Sinica Institute of Astronomy and Astrophysics and is funded by the Smithsonian Institution and the Academia Sinica.} of the 183 GHz maser emission from a region of low mass star formation, associated with the protostar SMM1. % ",
        "conclusions": "The conclusions of this letter are as follow \\begin{itemize} \\item The 183 GHz emission of Serpens SMM1 is masing from three distinct maser spots at 500-1200 AU, aligned in the direction of the red-shifted outflow, two of which are slightly blue-shifted. \\item The extended dust continuum emission detected with the SMA in Serpens SMM1 comes from a protostellar envelope with an estimated mass of 2.7 M$_\\odot$. \\item A resolved component detected on baselines up to $\\sim$400 k$\\lambda$ indicates the presence of disk with an estimated mass of 0.1 M$_\\odot$. \\item It is theorized that the maser spots originate in Kelvin-Helmholtz instabilities in the red-shifted outflow, created by the jet interactions with the envelope. Such interactions would temporarily create small regions of high density and temperature, providing the conditions for the 183 GHz transition to mase in low-mass protostellar environment. \\item No compact or extended thermal water emission is detected, in agreement with a model of a passively heated envelope. \\end{itemize} Research into the 183 GHz maser is promising as an additional constraint to jet mechanics and interaction of the jet with its surrounding material, both the envelope and the outflow.  In the next few years, the receivers of the SMA provide a unique opportunity to study the water maser emission at this wavelength in many low-mass YSOs.  In the long run, ALMA will be able to observe these lines at higher sensitivity and resolution, possibly probing the scales at which these masers emit. Combined with high-resolution VLA and VLBI observations of the 22 GHz maser, such observations can truly probe the physical structure of maser-emitting regions.  \\ack TvK is supported as an SMA postdoctoral fellow. Steve Longmore and Jes J{\\o}rgensen are thanked for discussion on fitting envelope models on continuum, and Elizabeth Humphreys for extensive discussion on maser excitation. TvK is also grateful to Lars Kristensen for providing the DUSTY model. The extensive ongoing discussions with Ewine van Dishoeck on many aspects of interstellar water are much appreciated.   \\begin{figure}[!ht] \\begin{center} \\includegraphics[width=400pt]{./figure1.ps} \\end{center} \\caption{Moment map of the three maser spots in SMM1. Maser spots are labeled with their respective number.  An X at the center 0,0 position marks the maximum intensity of the continuum (see Fig. \\ref{1:cont}), which coincides with the location of the protostar. The spectra at the three maser spots are shown . Due to the proximity of the spot 2 to both spot 1 and 3 ($\\sim$ 1 beamsize), the spectrum there is a blend of all three maser spots.  The beamsize is to 2.7$''$ by 4.0$''$ with a P.A. of 49$^\\circ$.} \\label{1:moment} \\end{figure}   \\begin{figure}[!ht] \\begin{center} \\includegraphics[width=220pt]{./figure2a.ps} \\includegraphics[width=220pt]{./figure2b.ps} \\includegraphics[width=230pt]{./figure2c.ps} \\includegraphics[width=230pt]{./figure2d.ps} \\end{center} \\caption{{\\it TOP:} The continuum map of SMM1 at 187 GHz.  Levels are in 2$\\sigma$, 4$\\sigma$, 6$\\sigma$,... with $\\sigma$=15 mJy/beam. {\\it BOTTOM:} The UV distance versus amplitude plot, with the mean visibility amplitude in that annulus on the vertical axis and baseline length on the horizontal. The triangles show the observed visibilities of SMM1 at 187 GHz. Errors shown are the standard deviation in the mean, while the solid line shows the expectation value assuming no signal. The dashed line shows the expected visibilities for an envelope model at 194 GHz (Kristensen et al. in prep). The crosses represent the residual visibilities if the envelope model is subtracted from the observed visibilities.} \\label{1:cont} \\end{figure}   \\begin{table}[!ht] \\caption{Properties of the maser spots} \\begin{center} \\begin{tabular}{l l l l l l} \\hline \\hline Maser spot &  FWHM  & $I_{\\rm{peak}}$  & $T_{\\rm{B}}$$^a$ &$V_{\\rm{LSR}}$ & Spat. Offset \\\\ & km s$^{-1}$ & Jy/beam & K &km s$^{-1}$ & RA,Dec ('','') \\\\ \\hline 1 & 0.7 & 463& 1871 & 9.2 & -3.0,0.2  \\\\ 2 & 1.5 & 243 & 982 &7.4 & -4.5,2.2\\\\ 3 & 1.3 & 328 & 1325 & 6.3 & -6.7,3.8\\\\ \\hline \\end{tabular} \\end{center} $^a$ assuming the emission fills the beam \\label{1:tab} \\end{table}"
    },
    "0910/0910.5625_arXiv.txt": {
        "abstract": "The decay of dark matter is predicted by many theoretical models and can produce observable contributions to the cosmic-ray fluxes. I shortly discuss the interpretation of the positron and electron excess as observed by PAMELA and Fermi LAT in terms of decaying dark matter, and I point out the implications for the Fermi LAT observations of the $\\gamma$-ray flux with emphasis on its dipole-like anisotropy. ",
        "introduction": "The most popular type of dark matter (DM) candidate, the weakly interacting massive particle (WIMP), can naturally reproduce the observed DM abundance due to effective self-annihilation in the early Universe, and today this same annihilation process could produce an observable contribution to the measured cosmic-ray fluxes on Earth. Such an indirect detection of DM is also possible if DM \\textit{decays} with a sufficiently large rate. There exist a number of interesting DM models (see \\textit{e.g.}~\\cite{models} and references therein) that predict the decay of DM on cosmological time scales, namely with lifetimes around and above $\\tau_\\text{DM}\\simeq \\mathcal{O}(10^{26}\\,\\text{s})$, which are typically required to be not in conflict with current observational limits. Among these models is the gravitino with a small violation of $R$-parity, motivated by requiring a consistent thermal history of the Universe, and the sterile neutrinos, whose long lifetime is due to tiny Yukawa couplings. The typical masses for these DM candidates lie in the $100\\,\\text{GeV}$ and the $10\\,\\text{keV}$ regime, respectively. Another interesting model with kinetically mixed hidden gauginos was also recently studied~\\cite{hg}. Even in models where DM is stable in the first place, the consideration of higher-dimensional operators often renders the DM particle unstable with cosmological lifetimes. Since the indirect detection signals from decay differ in general from the ones of annihilation, a dedicated study of decaying DM signals is mandatory. Below I will shortly review the $\\gamma$-ray and $e^\\pm$-signals that can come from DM decay, and I will discuss them in light of recent observations. ",
        "conclusions": "Many theoretical models predict the decay of DM on cosmological timescales, giving rise to an anomalous contribution to the observed cosmic-ray fluxes. The corresponding $\\gamma$-ray signals could show up as broad features over large angular distance in the $\\gamma$-ray sky. If decaying DM is the right explanation of the positron and electron excess observed by PAMELA and Fermi LAT, a corresponding $\\gamma$-ray signal with a large dipole-like anisotropy should be observed in the very near future with Fermi LAT. This anisotropy would be due to prompt radiation at high latitudes, and due to ICS radiation at lower latitudes, most prominent in a region of a few kpc around the galactic center. It is tempting to speculate that such an ICS signal already showed up in the Fermi LAT data, see Ref.~\\cite{Dobler:2009xz}."
    },
    "0910/0910.2233_arXiv.txt": {
        "abstract": "We have built a reliable and robust system that takes as input an astronomical image, and returns as output the pointing, scale, and orientation of that image (the astrometric calibration or WCS information).  The system requires no first guess, and works with the information in the image pixels alone; that is, the problem is a generalization of the ``lost in space'' problem in which nothing---not even the image scale---is known.  After robust source detection is performed in the input image, asterisms (sets of four or five stars) are geometrically hashed and compared to pre-indexed hashes to generate hypotheses about the astrometric calibration.  A hypothesis is only accepted as true if it passes a Bayesian decision theory test against a null hypothesis.  With indices built from the USNO-B Catalog and designed for uniformity of coverage and redundancy, the success rate is $>99.9~\\percent$ for contemporary near-ultraviolet and visual imaging survey data, with no false positives.  The failure rate is consistent with the incompleteness of the USNO-B Catalog; augmentation with indices built from the 2MASS Catalog brings the completeness to $100~\\percent$ with no false positives.  We are using this system to generate consistent and standards-compliant meta-data for digital and digitized imaging from plate repositories, automated observatories, individual scientific investigators, and hobbyists.  This is the first step in a program of making it possible to trust calibration meta-data for astronomical data of arbitrary provenance. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.2227_arXiv.txt": {
        "abstract": "We select 25,000 galaxies from the NEWFIRM Medium Band Survey (NMBS) to study the rest-frame $U~-~V$ color distribution of galaxies at $0<z\\lesssim2.5$.  The five unique NIR filters of the NMBS enable the precise measurement of photometric redshifts and rest-frame colors for 9,900 galaxies at $1<z<2.5$.  The rest-frame $U~-~V$ color distribution at all $z\\lesssim2.5$ is bimodal, with a red peak, a blue peak, and a population of galaxies in between (the green valley).  Model fits to the optical-NIR SEDs and the distribution of MIPS-detected galaxies indicate that the colors of galaxies in the green valley are determined largely by the amount of reddening by dust.  This result does not support the simplest interpretation of green valley objects as a transition from blue star-forming to red quiescent galaxies.  We show that correcting the rest-frame colors for dust reddening allows a remarkably clean separation between the red and blue sequences up to $z\\sim2.5$.   Our study confirms that dusty starburst galaxies can contribute a significant fraction to red sequence samples selected on the basis of a single rest-frame color (i.e. $U~-~V$), so extra care must be taken if samples of truly ``red and dead'' galaxies are desired.  Interestingly, of galaxies detected at \\24mum, 14\\% remain on the red sequence after applying the reddening correction. ",
        "introduction": "\\label{s:intro}  A powerful and observationally inexpensive tool for studying galaxy evolution is the color-magnitude or, more fundamentally, the color-mass diagram, in which the distribution of galaxies is bimodal with early-types residing on a well-defined sequence at red colors that is largely separate from a cloud of blue late-types \\markcite{blanton:03}(e.g. {Blanton} {et~al.} 2003).  Recent deep surveys have enabled the study of the evolution over $0<z<1$ of luminosity (\\markcite{bell:04}{Bell} {et~al.}, 2004, hereafter B04;~\\markcite{brown:07, faber:07}{Brown} {et~al.} 2007; {Faber} {et~al.} 2007) and mass \\markcite{borch:06}({Borch} {et~al.} 2006) functions divided between red and blue samples selected based on color-magnitude and color-mass diagrams.  A primary result of these studies is that the stellar mass in red galaxies builds up by a factor of $\\sim$2 between $z=1$ and $z=0$, while the stellar mass in blue galaxies remains roughly constant over the same redshift range.  This may imply a migration from the blue to red sequences (B04; \\markcite{faber:07}{Faber} {et~al.} 2007).  In the simplest models the evolution from blue to red colors is largely driven by simple aging of stellar populations, with the rate of objects entering the red sequence modulated by the particular feedback processes at work that regulate star formation \\markcite{schawinski:07}(e.g., {Schawinski} {et~al.} 2007).  Recently, several surveys have started to push the systematic study of galaxy colors to $z>1$.  There is some evidence that the color bimodality persists up to $z=2$, though the contrast separating the red and blue sequences is low.  This is either because spectroscopic samples that span a broad range of colors are small \\markcite{giallongo:05,franzetti:07, cassata:08, kriek:08}({Giallongo} {et~al.} 2005; {Franzetti} {et~al.} 2007; {Cassata} {et~al.} 2008; {Kriek} {et~al.} 2008) or because larger photometric samples suffer from uncertain photometric redshifts at $z>1.5$, which leads to uncertain rest-frame colors \\markcite{taylor:09a}({Taylor} {et~al.} 2009a).  \\markcite{wuyts:07}{Wuyts} {et~al.} (2007) (hereafter W07) and \\markcite{williams:09}{Williams} {et~al.} (2009) (hereafter W09) demonstrate that quiescent galaxies up to $z\\sim2$ can be fairly cleanly separated from the actively star-forming population when multiple rest-frame colors are used.  In this Letter we use the NEWFIRM Medium Band Survey \\markcite{nmbs}(NMBS; {van Dokkum} {et~al.} 2009) to explore the evolution of the galaxy color bimodality over $0<z\\lesssim2.5$.  We show that two independent methods for accounting for dust in objects with intermediate colors allows the separation of a clearly-defined red sequence up to at least $z\\sim2.5$.  We assume a $\\Lambda$CDM cosmology with $\\Omega_M=0.3$, $\\Omega_\\Lambda=0.7$, and $H_0 = 70\\ \\mathrm{km\\ s}^{-1}\\ \\mathrm{Mpc}^{-1}$ throughout.  Broad-band magnitudes and colors are given in the AB system.  \\begin{figure} \\plotone{f1.eps} \\figcaption{Data quality of the NMBS.  $u\\mbox{--}8\\mu\\mathrm{m}$ SEDs are shown for six objects at $1.7 < z < 2$ that have $K=22\\pm0.1$, the median magnitude in this redshift range.  From top to bottom, the objects shown have $A_V = \\left[0.0,0.6,0.9,1.5,2.1,2.5\\right]$ (\\S\\ref{s:data}). The arrows point towards the measured MIPS fluxes and show that the objects whose SED fits indicate large amounts of dust are also detected at \\24mum. \\label{f:seds}} \\end{figure}  \\begin{figure*} \\plotone{f2.eps} \\figcaption{Rest-frame $\\left(U~-~V\\right)^\\prime$ color as a function of redshift for objects in both NMBS fields with $K < 22.8$.  The color plotted is corrected for the slope of the (local) color-magnitude relation (Eq. \\ref{eq:cmr}).  The density of points is shown as a 2D histogram scaled so that the maximum bin in each column is shaded black.  The solid diagonal line in each panel shows the evolution of the CMR zeropoint found by B04.  Objects detected and not detected with MIPS at \\24mum\\ are shown in the center and right panels, respectively.  Removing the MIPS-detected objects allows a remarkably clean separation of the red and blue galaxy sequences.  The hatched regions indicate where our $K<22.8$ sample is less than 90\\% complete at $M_\\star=10^{10.7}M_\\odot$.  The curvature of the completeness line is due to the observed $K$ magnitude selection.  As discussed extensively by, e.g., \\markcite{taylor:09a}{Taylor} {et~al.} (2009a), $K$-selected samples are more complete for blue galaxies than for red galaxies at a fixed mass. \\label{f:uv_z_av}} \\end{figure*} ",
        "conclusions": "\\label{s:discussion}  We have shown that either removing MIPS-detected galaxies or applying a reddening correction to rest-frame $U~-~V$ colors allows us to cleanly divide the population of galaxies with red rest-frame colors among dusty-starburst and ``red and dead'' galaxies.  The latter, which have \\textit{intrinsically} red colors characteristic of evolved stellar populations, form what we call the ``dead sequence'', which is in place at all $0\\lesssim z\\lesssim2.5$.  This term was used previously by \\markcite{romeo:08}{Romeo} {et~al.} (2008), who consider the sub-sample of red-sequence galaxies with $\\log \\left(\\mathrm{sSFR\\cdot yr}\\right) < -11$;  Figure \\ref{f:cmd} shows that the cloud of galaxies with red dust-corrected colors do, in fact, have very low sSFR.  We note, however, that there could be galaxies on what we consider the ``dead sequence'' with somewhat higher sSFR than in \\markcite{romeo:08}{Romeo} {et~al.} (2008).  Obtaining these results at $z\\gtrsim1.5$ is made possible by the NEWFIRM Medium-band Survey, which provides high-quality \\zphot\\ and well-sampled SEDs of red galaxies at redshifts beyond the current limits of large spectroscopic surveys \\markcite{nmbs}({van Dokkum} {et~al.} 2009): of the $\\sim$3000 NMBS objects with measured spectroscopic redshifts, only 9 have $U-V>1.4$ at $z>1.2$.  Conversely, there are 2826 objects that satisfy these criteria in the full photometric sample.  The results presented here are consistent with those of W09, who show that quiescent galaxies can be separated from dusty star-forming galaxies in the $U~-~V$ vs. $V~-~J$ color-color diagram out to $z=2$.  The reddening-corrected dead sequence galaxies discussed here fall nicely within the $UVJ$ quiescent selection criteria, though the reddening correction is critical for detecting the bimodal color distribution in a color-mass diagram at $z>1.5$, explaining why it is not seen by W09.  We find the color evolution of dead sequence galaxies to be consistent with that found by B04, who show that it can be roughly explained by simple stellar populations formed at $z>2$ that age passively to the present day.  We show, however, that additional care must be taken when defining red sequence galaxy samples as done by, e.g., B04.  The bottom panel of Figure \\ref{f:average_av} shows that a red galaxy selection based on a single rest-frame color will contain a significant fraction of dusty-starburst galaxies ($\\sim$20\\% with $A_V>1$) that would otherwise be much bluer (see also \\markcite{cimatti:02}{Cimatti} {et~al.} 2002; \\markcite{brammer:07}{Brammer} \\& {van Dokkum} 2007; W07; \\markcite{gallazzi:09}{Gallazzi} {et~al.} 2009; W09; \\markcite{wolf:09}{Wolf} {et~al.} 2009).  The reddening correction dramatically reduces the number of objects in the green valley, so such a correction would be necessary when using those objects to estimate the rate at which galaxies move from the blue to red sequences \\markcite{martin:07}({Martin} {et~al.} 2007).  For example, a sparsely-populated green valley at most redshifts could constrain the rate at which star-forming galaxies are quenched, which in turn may help constrain the processes that regulate star formation in theoretical models.  As might be expected, the large majority of objects detected at \\24mum\\ have reddening-corrected colors that place them on the blue-sequence \\markcite{reddy:06}({Reddy} {et~al.} 2006, see also Figure 3). A number of MIPS-detected galaxies, however, remain on the red sequence even after applying the reddening correction:  roughly one in seven MIPS objects have dust-corrected colors within $0.5$ mag of the red sequence ridgeline drawn in Figure \\ref{f:uv_z_av}.  The dead sequence galaxies detected with MIPS have de-reddened colors that are slightly bluer than those without MIPS detections (Figure \\ref{f:cmd}).  This could perhaps be due to low levels of recent star formation or to AGN contamination of the UV-optical light, though this tentative result requires more detailed investigation because, for example, it is not clear that the dust correction or even the extinction law itself are appropriate for these galaxies.  Figures \\ref{f:uv_z_av} and \\ref{f:cmd} demonstrate that the dead sequence persists to $z=2\\mbox{--}2.5$, which is consistent with the discovery of a small (spectroscopic) sample of red sequence objects with low sSFR at $z\\sim2.3$ described by \\markcite{kriek:08}{Kriek} {et~al.} (2008).  Such studies of red sequence galaxies found among the field are complementary to others that target the higher density environments of clusters at $z\\gtrsim1$ \\markcite{mei:09}(e.g., {Mei} {et~al.} 2009) and proto-clusters at $z\\sim2$ \\markcite{zirm:08}({Zirm} {et~al.} 2008). Mechanisms that have been proposed to suppress star formation in massive galaxies, including the much-discussed feedback from AGN and gravitational heating from cosmological accretion \\markcite{naab:07,dekel:08}({Naab} {et~al.} 2007; {Dekel} \\& {Birnboim} 2008), must be able to explain the formation of the dead sequence in a variety of environments as early as $z\\sim2.5$. Finally, we note that quantitative constraints on the assembly and evolution of passively-evolving galaxies require studies of the evolution of their mass function, possibly combined with structural information (see, e.g., B04; \\markcite{faber:07}{Faber} {et~al.} 2007; \\markcite{cimatti:08}{Cimatti} {et~al.} 2008; \\markcite{vandokkum:08}{van Dokkum} {et~al.} 2008, and many other studies).  The NMBS is ideally suited to help disentangle these processes by providing accurate photometric redshifts and colors over a large area in the redshift regime of the most rapid build-up of massive galaxies."
    },
    "0910/0910.0014_arXiv.txt": {
        "abstract": "Finding electromagnetic (EM) counterparts of future gravitational wave (GW) sources would bring rich scientific benefits.  A promising possibility, in the case of the coalescence of a super-massive black hole binary (SMBHB), is that prompt emission from merger-induced disturbances in a supersonic circumbinary disk may be detectable. We follow the post--merger evolution of a thin, zero-viscosity circumbinary gas disk with two-dimensional simulations, using the hydrodynamic code FLASH.  We analyze perturbations arising from the 530 ${\\rm km~s^{-1}}$ recoil of a $10^6 M_\\odot$ binary, oriented in the plane of the disk, assuming either an adiabatic or a pseudo--isothermal equation of state for the gas.  We find that a single-armed spiral shock wave forms and propagates outward, sweeping up $\\sim 20\\%$ of the mass of the disk. The morphology and evolution of the perturbations agrees well with those of caustics predicted to occur in a collisionless disk. Assuming that the disk radiates nearly instantaneously to maintain a constant temperature, we estimate the amount of dissipation and corresponding post-merger light--curve.  The luminosity rises steadily on the time--scale of months, and reaches few $\\times 10^{43}$ erg/s, corresponding to $\\approx 10\\%$ of the Eddington luminosity of the central SMBHB.  We also analyze the case in which gravitational wave emission results in a $5\\%$ mass loss in the merger remnant.  The mass-loss reduces the shock overdensities and the overall luminosity of the disk by $\\approx 15-20\\%$, without any other major effects on the spiral shock pattern. ",
        "introduction": "The gravitational waves (GWs) produced during the late stages of the merger between super-massive black holes (SMBHs), with masses of $\\sim (10^4$--$10^7)\\,{\\rm M_\\odot}/(1+z)$, out to redshifts beyond $z\\approx 10$, are expected to be detectable in the next decade by the {\\it Laser Interferometric Space Antenna} ({\\it LISA}) satellite. While the GW signatures themselves will be a rich source of information, identifying the electromagnetic (EM) counterpart of the LISA source would open up a whole new range of scientific opportunities, from black hole astrophysics to fundamental aspects of gravitational physics and cosmology \\citep{koc07,HKM09a,Phinney,Bloom}.  Whether the EM counterparts of the SMBHBs detected by {\\it LISA} can be uniquely identified depends on the accuracy of localization by {\\it LISA}, and on the nature of the EM emission.  For the typical SMBHB detected by {\\it LISA} at $z\\approx 1-2$, the sky position and the redshift will be known to the accuracy of $\\delta(\\Delta\\Omega)\\sim$ few $\\times$ 0.1 square degrees, and $\\delta z\\approx$ few $\\times$ 0.01, respectively (the latter limited by weak lensing errors; e.g. \\citealt{hh05,koc06}).  Within this three--dimensional error volume, there will be of order $\\sim 100$ candidate galaxies, at the optical magnitude limit expected to host such SMBHBs \\citep{koc08}.  If the coalescing SMBHs themselves produce bright emission comparable to luminous quasars, or are associated with some other, similarly rare subset comprising $\\lsim 1\\%$ of all galaxies (such as ultra--luminous infrared galaxies), then a unique counterpart may be identified among these candidates \\citep{koc06}.  However, having a prediction for the spectrum and the light--curve of a coalescing SMBHB -- and therefore knowing what characteristic signatures to look for -- would make such identifications both more likely and more reliable.  A promising possibility is that the coalescing SMBHB produces a {\\em variable or transient} EM signal, which can be uncovered by suitably designed EM observations, either concurrently with or following the {\\it LISA} detection \\citep{koc08,lh08,HKM09a}.  The dense nuclear gas around the BH binary is expected to cool rapidly, and settle into a rotationally supported, thin circumbinary disk \\citep[e.g.][]{barnes,elcm05}.  The evolution of a SMBHB embedded in such a disk has been studied in various idealized configurations \\citep[e.g.][]{an02,liu03,mp05,dot06,mm08,hayasaki09,cuadra09,HKM09b}. Generically, at small orbital separations when the binary is detectable by {\\it LISA}, the orbit decays rapidly due to GW emission. If the disk is thin, the binary torques create a central cavity, nearly devoid of gas, within a region about twice the orbital separation \\citep[e.g.][]{al94}.  Whether the decaying SMBHBs produce bright emission during this stage is not well understood.  If the central cavity were truly empty, no gas would reach the SMBHBs, and any emission produced farther out in the disk would likely be weak. On the other hand, numerical simulations suggest residual gas inflow into the cavity \\citep{al96,mm08,hayasaki07,hayasaki08,cuadra09}, which may plausibly accrete onto the BHs, producing non--negligible EM emission.  A different possibility, and the focus of the present paper, is that variable EM signatures are produced in the gas disk {\\it promptly after} the coalescence of the SMBHB.  The burst of GWs emitted during the last stages of the coalescence results in a corresponding, nearly instantaneous reduction in the binary's rest mass. Furthermore, when compact objects coalesce asymmetrically, the linear momentum carried away by the GWs imparts a ``kick'' to the center of mass of the system \\citep{Bekenstein, Fitchett}.  Recent break--through in numerical relativity has allowed accurate calculations of these effects, showing that the mass loss is typically several percent \\citep[e.g.][and references therein]{tm08}, while kick velocities are typically several hundred ${\\rm km~s^{-1}}$, but can be as high as 4,000 ${\\rm km~s^{-1}}$ \\citep{baker06,baker07,baker08,campanelli07a,campanelli07b,gonzalez07a,gonzalez07b, herr07a,herr07b,herr07c,koppitz07}.  The circumbinary gas will respond promptly (on the local orbital timescale) to such dynamical disturbances.  Importantly, the gaseous disk outside the inner cavity is expected to cool efficiently and become geometrically thin, implying that the orbital motion of the gas is supersonic.  This gas is therefore susceptible to prompt shocks, which could, in principle, produce a detectable transient EM signature \\citep{mp05}.  \\citet{Lippai} considered the motion of collisionless test particles in such a disturbed disk.  As long as the particles remain bound to the central SMBHB, they follow elliptical Kepler orbits (in the inertial frame of the SMBHB).  These orbits cross, and produce a characteristic, outward--propagating spiral caustic.  While these results were obtained in a pressureless ``dark matter'' disk, they suggest that shocks, with a similar pattern, will indeed arise in the gas, on times--scales of $\\sim$weeks to $\\sim$months after the coalescence.  Similar conclusions were reached by \\citet{Shields} and by \\citet{Schnittman}, using N-body simulations and semi--analytical arguments, respectively, to show that a bright, post--merger ``flare'' may occur.  Both of these studies focused on the evolution of disks around more massive BHs on longer ($\\sim 10^4$yr) time--scales, and proposed detecting the flare by monitoring a population of active galactic nuclei (AGN).  By comparison, our analysis here focuses on the disks around lower--mass BHs and on the shorter time--scales of several weeks to a year, which are relevant for follow--up observations triggered by {\\it LISA} detections.  Motivated by the above, in the present paper, we follow up on earlier work, and compute the response of the gas disk including the effects of gas pressure using hydrodynamical simulations.  Our main goals are (i) to assess shock formation in hydrodynamical disks, and (ii) to estimate the amount of dissipation and the resulting light--curve produced by the disk.  In particular, for the latter, we will use simulations with a pseudo--isothermal equation of state. This corresponds implicitly to strong dissipation, and should represent an approximate upper limit on the disk luminosity.  Our computations are performed with the publicly available code FLASH \\citep{Fryxell}.  Two other, independent studies have recently used 3--dimensional simulations to address the response of gas disks to mass--loss \\citep{ONeill} and to both mass--loss and kicks \\citep{Megevand}.  The main difference between our analysis and these previous works is our inclusion of runs with a pseudo--isothermal equation of state, which modifies, qualitatively, the expected EM signature.  In particular, \\citet{ONeill} and \\citet{Megevand} both run simulations with an adiabatic equation of state, and find that in most cases, the gas disk typically {\\em dims} following the merger. We observe a similar effect in our adiabatic runs, and attribute it to an overall dilution of the gas density.  However, the gas develops significantly larger density contrasts in our pseudo--isothermal runs, and we argue that this can result in a significant {\\it brightening} of the system.  Another important difference between our study and those of \\citet{ONeill} and \\citet{Megevand} is that we perform two--dimensional simulations, whereas \\citet{ONeill} and \\citet{Megevand} both utilized 3D simulations.  While we have sacrificed resolving the vertical disk structure, this allows us to simulate a disk that extends to much larger radii ($10^4$ Schwarzschild radii), and follow the entire disk for a much longer period ($\\sim1$ yr for a $10^6~{\\rm M_\\odot}$ binary), in order to cover the regime relevant for follow-up observations of {\\it LISA} events.  Given the expected size of the inner cavity in the disk, $\\sim 100 R_s$, studying these outer regions of the disk, and following the disk evolution on the correspondingly longer time--scale, is especially important. Further differences between our study and those of \\citet{ONeill} and \\citet{Megevand} will be discussed in \\S~\\ref{sec:Others} below.  The rest of this paper is organized as follows. In \\S~\\ref{sec:ProblemSetup}, we describe our computational setup, including details of the simulations, and initial conditions. In \\S~\\ref{sec:Results}, we present and discuss our main results. These include the impact of the kick (\\S~\\ref{sec:PostMerger}) and the mass--loss (as well as both effects in combination; \\S~\\ref{sec:MassLoss}), using constant surface density disks.  In \\S~\\ref{sec:Observational}, we discuss our estimate of the post--merger light curves.  In \\S~\\ref{sec:Alphadisks}, we repeat our calculations at higher resolution and with more realistic initial gas density and temperature profiles (adapted from a standard $\\alpha$-disk).  In \\S~\\ref{sec:Numerical}, we discuss possible numerical issues, and in \\S~\\ref{sec:Others}, we compare our results to previous work. Finally, in \\S~\\ref{sec:Conclude}, we summarize our main conclusions and their implications, and outline natural future extensions of this study. ",
        "conclusions": "\\label{sec:Conclude}  In this paper, we used two--dimensional hydrodynamical simulations of a circumbinary disk, to follow the effects of a velocity recoil and mass-loss of the central black hole binary, following the merger of the two black holes. From a suite of runs, we are able to draw two basic conclusions.  First, the outward--propagating spiral shocks that develop in our simulations follow a pattern very similar to the caustics identified in collisionless disk \\citep{Lippai}.  Gas pressure has a modest overall impact on the propagation and on the overall morphology of the shocks.  On the other hand, we find that the gas pressure, and the assumed equation of state, has a strong impact on the overdensities that develop and on the strengths of the shocks: isothermal, low--temperature disks have larger overdensities and stronger shocks.  Second, we have estimated an upper limit on the luminosity emerging from the disk experiencing a BH kick with $v_{\\rm kick}=530~{\\rm km~s^{-1}}$, by measuring the effective dissipation that occurs, implicitly, in the simulations when an isothermal condition is imposed on the gas.  The resulting luminosity is of order $10\\%$ of the Eddington limit for the $10^6 M_\\odot$ BH used in the simulation, which suggests that the after--glow may be bright enough to be detectable.  If the pre-merger disk has a luminosity below that of a standard steady-state thin accretion disk (due to the evacuation of the inner regions of the disk), then the merger-induced kick will cause a significant (order-of-magnitude, or larger) brightening of the disk.  We also estimated the effective black--body temperature of the radiation emerging from the optically thick disk. We found that as the disk brightens, the characteristic frequency increases with time, possibly offering a unique signature of the kick--induced emission. When the accompanying mass--loss of the merger remnant is included in our simulations, the density contrasts and luminosity from the spiral shocks decrease somewhat, but do not change dramatically in overall behavior.  While our results are encouraging, and suggest that the EM signature of the disturbed circumbinary disk may be detectable, this conclusion has to be verified in the future in improved models. In particular, a better estimate of the thermodynamics of the disk, with a realistic treatment of radiative cooling, should reveal how close the luminosity is to the values we obtained here, using the isothermal assumption; by using three--dimensional simulations that resolve the vertical disk structure, and by following the vertical transfer of the radiation produced should clarify the robustness of our conclusions about the evolution of the characteristic emergent frequency."
    },
    "0910/0910.0687_arXiv.txt": {
        "abstract": "Emission of high energy (HE) photons above 100 MeV that is delayed and lasts much longer than the prompt MeV emission has been detected from several long duration gamma ray bursts (LGRBs) and short hard bursts (SHBs) by the Compton, Fermi and AGILE gamma ray observatories. In this paper we show that the main observed properties of this HE emission are those predicted by the cannonball (CB) model of GRBs: In the CB model all the observed radiations in a GRB are produced by the interaction of a highly relativistic jet of plasmoids (CBs) with the environment. The prompt X-ray and MeV $\\gamma$-ray pulses are produced by inverse Compton scattering (ICS) of glory photons -photons scattered/emitted into a cavity created by the wind/ejecta blown from the progenitor star or a companion star long before the GRB- by the thermal electrons in the CBs. A simultaneous optical and high energy emission begins shortly after each MeV pulse when the CB collides with the wind/ejecta, and continues during the deceleration of the CB in the interstellar medium. The optical emission is dominated by synchrotron radiation (SR) from the swept-in and knocked-on electrons which are Fermi accelerated to high energies by the turbulent magnetic fields in the CBs, while ICS of these SR photons dominates the emission of HE photons. The lightcurves of the optical and HE emissions have approximately the same temporal behaviour but have different power-law spectra. The emission of very high energy (VHE) photons above 100 TeV is dominated by the decay of $\\pi^0$'s produced in hadronic collisions of Fermi accelerated protons in the CBs. The CB model explains well all the observed radiations, including the high energy radiation from both LGRBs and SHBs as demonstrated here for GRB 090902B and SHB 090510. ",
        "introduction": "During nearly 20 years after the launch of the Compton Gamma Ray Observatory (CGRO), the Burst And Transient Source Experiment (BATSE) on board the CGRO has detected and measured light curves and spectra (Kaneko et al.~2008)  in the sub-MeV range of several thousands gamma ray bursts (GRBs). Higher-energy observations with the EGRET instrument aboard CGRO were limited to those GRBs which happened to be in its narrower field of view. Its large calorimeter measured the light-curves and spectra of several GRBs in the 1-200 MeV energy range. Seven GRBs were detected also with the EGRET spark chamber, sensitive in the 30 MeV - 10 GeV energy range. The EGRET detections indicated that the spectrum of bright GRBs extends beyond 1 GeV (Hurley et al.~1994) with no evidence for a spectral cut-off (see, e.g., Dingus~1995,~2001 and references therein). However, a few GRBs, such as 940217 (Hurley et al.~1994) and 941017 (Gonzalez et al.~2003), showed evidence that the high energy component has a slower temporal decay than that of the sub-MeV emission, suggesting that, at least in some cases, it is not a simple extension of the main component, but originates from a different emission mechanism and/or region. This has been confirmed recently by observations of high energy photons in the 30 MeV - 300 GeV range from several GRBs with the AGILE (GRB 080514B: Giuliani et al.~2008, GRB 090510: Giuliani et al.~2009) and the Fermi large area telescope (LAT) (e.g., GRB 080916C: Abdo et al.~2009a, GRB 090902B: Bissaldi et al.~2009 and GRB 090510: Ghirlanda et al.~2009). The arrival times of the high energy photons did not coincide with the times of the brightest peaks seen at hard X-rays and MeV $\\gamma$-rays. Also the high energy emission lasted much longer time than that of the prompt keV-MeV emission.    The detection of higher energy gamma rays is affected by pair production in their collisions with the extragalactic infrared background light (Nikishov 1961) resulting in an absorption which is a strong function of redshift and energy. Recent estimates (Primack et al.~2005), validated by HESS observations (Aharonian et al.~2006) predict an optical depth of roughly unity to 500 GeV photons emitted at a redshift $z\\!=\\!0.2$ and to 10 TeV photons at $z=0.05$. The average redshifts of LGRBs and SHBs are much larger, $z\\!=\\!2.2$ and $z\\!=\\!0.5$, respectively. Despite of the strong attenuation of high energy gamma rays in the intergalactic medium (IGM), there have been several claims in the past of detections at the 3 sigma level of TeV gamma-rays from GRBs (see, e.g., Atkins et al.~2003 and references therein). However, more recently, no GRB was conclusively detected in the range 100 GeV to 100 TeV by the ground based water Cherenkov detector Milagro and by the air Cherenkov telescopes MAGIC, Whipple, HESS and VERITAS. Moreover in all previous cases of reported detection of TeV gamma-rays from a GRB, the GRB redshift was not known. If their redshifts are similar to those of ordinary GRBs then their TeV gamma-rays are strongly absorbed by the extragalactic background light (EBL), implying that the TeV emission if detected must be extraordinarily energetic, i.e., with a much larger fluence than that emitted in X-rays and MeV $\\gamma$-rays.   Most theoretical models of high energy photon emission in GRBs relied on the standard fireball models of GRBs (for recent reviews see, e.g., Piran~2005; M\\'{e}sz\\'{a}ros~2006; Zhang~2007). In these models, synchrotron radiation of electrons accelerated by `internal shocks' in collisions between conical shells ejected by the GRB's central engine produces the prompt GRB emission and a blast wave (external shock) driven into the circumburst medium generates their afterglow.  The high energy radiation was suggested to be produced either by inverse Compton scattering of the synchrotron radiation (the so called `synchrotron self Compton mechanism') by the shock-accelerated electrons (e.g., Dermer et al.~2000), by the decay of $\\pi^0$ photo produced in collisions of shock accelerated hadrons with synchrotron photons in the expanding fireball (Waxman \\& Bahcall~1997), or by synchrotron radiation from ultra high energy protons (Totani~1998) allegedly accelerated in the GRB fireball (Waxman~1995; Milgrom \\& Usov~1995; Vietri~1995). All these models that were based on the standard fireball model of GRBs predict simultaneous emissions at all energy bands. This is in conflict with the observed delayed emission of the high energy photons that lasts much longer.  In the cannonball (CB) model (see e.g. Dado, Dar \\& De R\\'ujula, hereafter DDD, 2009a,b and references therein), GRBs and their afterglows (AGs) are produced by bipolar jets of highly relativistic plasmoids (CBs) ejected in violent stellar processes. The prompt MeV $\\gamma$-rays and hard X-rays are produced by inverse Compton scattering (ICS) of glory photons - photons emitted/scattered into a cavity formed by the wind/ejecta puffed by the progenitor or a companion star long before the GRB. Synchrotron radiation (SR) emitted from the electrons of the ionized wind/ejecta that are swept into the CBs and are Fermi accelerated by their turbulent magnetic fields dominates their `prompt' optical emission (e.g., DDD2009a and references therein) which begins when the CBs reach the wind/ejecta. In this paper we show that in the CB model ICS of this SR (e.g., Dado \\& Dar~2005) dominates the HE emission from long GRBs and SHBs. This HE emission begins simultaneously with the `prompt' optical emission that lags after the prompt X-ray and MeV $\\gamma$-ray emission and lasts much longer, has the same lightcurve as the optical emission but with a power-law spectrum that is identical to that of the X-ray emission. It describes well the HE observations, as demonstrated in the paper for GRB 090902B and SHB 090510.    Because of the Klein-Nishina effect, $E^2\\,dn/dE$ of ICS of the prompt SR in GRBs peaks near 10 TeV and is cutoff at a few tens of TeV when the rate of synchrotron energy losses by electrons in the CBs exceeds the rate of energy gain by Fermi acceleration.  The prompt decay of $\\pi^0$'s produced in the collisions between the Fermi accelerated nuclei and the ambient matter in the CBs produces a power-law spectrum that extends to much higher energies where it dominates the HE emission. However, like in blazars, the observed flux of TeV photons from distant GRBs and SHBs is stronglyly suppressed by pair production in collisions with the extragalactic background photons and only relatively nearby GRBs and SHBs might be detected in TeV photons by the large ground based HE gamma ray telescopes such as HESS, MAGIC and VERITAS. ",
        "conclusions": "In the cannonball (CB) model, high energy emission from GRBs and SHBs is a natural consequence of the model. The dominant leptonic and hadronic emission mechanisms are ICS of SR by Fermi accelerated electrons in the collision of the jet of highly relativistic CBs with the wind/ejecta blown from the progenitor or companion star long before the GRB, and the decay of $\\pi^0$ produced in hadronic collisions between the CB nuclei and the nuclei of the hadronic matter (ejecta/wind/ISM) that the jet passes through. The main predictions of the model  for the early-time high energy emission are: \\begin{itemize} \\item{} Each prompt keV-MeV  pulse is followed by a delayed HE emission which lasts much longer.  \\item{} The HE emission coincides in time with the optical emission.  \\item{} The light curve of the HE emission is roughly proportional to that of the unextinct optical emission. The decay of both the `prompt' optical and the HE emissions is a power-law with an index $\\alpha$=1+$\\beta_0\\!\\sim 1.5$.  \\item{} The spectrum of the HE component is a simple power-law with a spectral index approximately equal to that of the X-ray afterglow, i.e., $\\Gamma_{LAT}\\!\\approx\\!\\Gamma_{X}$.  \\item{} The HE emission extends to very high energies, where the observed radiation is strongly attenuated by $e^+e^-$ pair productin in the EBL.   \\item{} The equivalent isotropic energy of the HE component in bright GRBs can reach, and even exceed that in the sub-MeV range.  \\item{} The neutrino counterpart of the hadronic emission of HE $\\gamma$-rays from the most luminous GRBs is barely detectable in $km^3$ underwater/under-ice Cherenkov neutrino telescopes (DD2008).  \\end{itemize} \\noindent These predictions are consistent with the present available data on high energy emission from GRBs obtained from the gamma ray satellites and from the large air, ground and underground Cherenkov telescopes. In particular, the main observed properties of the HE emission measured with the Compton, Fermi and AGILE gamma ray satellites are well reproduced by the CB model as demonstrated in this paper for GRB 090902B and SHB 090510.   An observational proof of the ICS origin of the HE gamma ray emission from GRBs requires simultaneous detections of the prompt optical and HE emissions. It is highly desireable that the LAT team improves their automatic analysis, in order to deliver a GRB position within a few seconds after the LAT detection. Even a few tens of seconds will be very useful. This will provide a stringent test of models of HE emission from GRBs (and from other HE transient sources such as blazars, microquasars, pulsars, etc) and help pin down their production mechanism. Extremely optically-luminous GRBs, such as GRB 080319 where the prompt optical emission was resolved into individual flares (Racusin et al.~2008), may show that also the prompt high energy emission consists of HE flares which are associated with and follow promptly each individual keV-MeV pulse, as expected in the CB model.  \\newpage"
    },
    "0910/0910.0478_arXiv.txt": {
        "abstract": "We construct one-zone steady-state models of cosmic ray (CR) injection, cooling, and escape over the entire dynamic range of the FIR-radio correlation (FRC), from normal galaxies to starbursts, over the redshift interval $0 \\le z \\le 10$.  Normal galaxies with low star-formation rates become radio-faint at high $z$, because Inverse Compton (IC) losses off the CMB cool CR electrons and positrons rapidly, suppressing their nonthermal radio emission.  However, we find that this effect occurs at higher redshifts than previously expected, because escape, bremsstrahlung, ionization, and starlight IC losses act to counter this effect and preserve the radio luminosity of galaxies.  The radio dimming of star-forming galaxies at high $z$ is not just a simple competition between magnetic field energy density and the CMB energy density; the CMB must also compete with every other loss process.  We predict relations for the critical redshift when radio emission is significantly suppressed compared to the $z \\approx 0$ FRC as a function of star-formation rate per unit area.  For example, a Milky Way-like spiral becomes radio-faint at $z \\approx 2$, while an M82-like starburst does not become radio-faint until $z \\approx 10 - 20$.  We show that the ``buffering'' effect of non-synchrotron losses improves the detectability of star-forming galaxies in synchrotron radio emission with EVLA and SKA. Additionally, we provide a quantitative explanation for the relative radio brightness of some high-$z$ submillimeter galaxies.  We show that at fixed star formation rate surface density, galaxies with larger CR scale heights are radio bright with respect to the FRC, because of weaker bremsstrahlung and ionization losses compared to compact starbursts.  We predict that these ``puffy starbursts'' should have steeper radio spectra than compact galaxies with the same star-formation rate surface density.  We find that radio bright submillimeter galaxies alone cannot explain the excess radio emission reported by ARCADE2, but they may significantly enhance the diffuse radio background with respect to a naive application of the $z \\approx 0$ FRC. ",
        "introduction": "The radio and FIR luminosities of star-forming galaxies are linearly correlated over three decades in luminosity \\citep[][]{vanDerKruit71,vanDerKruit73,deJong85,Helou85,Condon92,Yun01}.  This FIR-radio correlation (FRC) is driven by star-formation.  Starlight emitted by young, massive stars in the optical and UV is reprocessed by dust into far-infrared radiation.  At the same time, the supernova remnants of massive stars accelerate cosmic rays (CRs), which produce synchrotron radio emission in the galaxies' magnetic fields.  The relation at $z \\approx 0$ is linear except in the very lowest luminosity galaxies \\citep{Condon91b,Yun01,Bell03}.  Since the FRC has less than a factor of 2 scatter in the local Universe, radio luminosity is often used as a tracer of star-formation at both low and high redshift \\citep[e.g.,][]{Cram98,Mobasher99}.  Although the FRC has been mainly studied in the low-$z$ universe, there has been recent interest in understanding how it applies to star-forming galaxies at high $z$.  There have been several conflicting results observationally.  A number of studies have found that the FRC holds unchanged or with little evolution at high redshifts (e.g., \\citealt{Appleton04,Ibar08,Rieke09,Murphy09,Younger09,Garn09,Sargent10,Ivison10,Younger10}; \\citealt{Persic07} suggest radio-dim ULIRGs at high $z$).  Submillimeter galaxies, however, seem to be radio-bright, by a factor of $\\sim 3$ (\\citealt{Kovacs06,Vlahakis07,Sajina08,Murphy09,Seymour09,Murphy09c,Michalowski09,Michalowski10}).  In particular, \\citet{Murphy09} and \\citet{Michalowski09} argue that their samples of submillimeter galaxies are intrinsically radio bright with respect to the local FRC, with no significant contamination from radio bright AGN.  In addition to the work on individual star-forming galaxies, there have also been new investigations into both the recently-resolved FIR and the unresolved GHz radio backgrounds.  The FIR background was detected by COBE and has long been attributed to star-formation \\citep[see the reviews by][]{Hauser01,Lagache05}.  The star-formation origin of the FIR background was recently confirmed by BLAST \\citep{Devlin09,Pascale09}.  A detection of the extragalactic GHz radio background has also been recently reported by ARCADE2 \\citep{Fixsen09,Seiffert09}.  Predictions for the strength of the diffuse radio background from star formation rely on the FRC holding out to $z \\approx 1 - 2$ \\citep[e.g.,][]{Haarsma98,Dwek02}, where most star-formation occurs.  In this paper, we describe how and why the linearity and normalization of the FRC are broken (or preserved) for different types of galaxies at high $z$.  We describe the basic physics of our model and the elements of the high-$z$ universe that physically allow the FRC to break down in \\S~\\ref{sec:Theory}.  In \\S~\\ref{sec:Procedure}, we briefly describe our underlying procedure.  In \\S~\\ref{sec:Results}, we discuss our results, providing predictions for when star-forming galaxies should be more or less radio luminous than the $z \\approx 0$ FRC suggests.   We also summarize the implications of our predictions for radio emission as a star formation tracer in \\S~\\ref{sec:SFTracer} and estimate the effects of radio-bright SMGs on the diffuse radio background in \\S~\\ref{sec:RadioBackground}.  Throughout this paper, primed quantities denote the source's rest-frame and unprimed quantities are for the observer-frame. ",
        "conclusions": "\\label{sec:Results} \\label{sec:Discussion}  \\subsection{The Evolving FIR-Radio Correlation} \\label{sec:FRCEvolution}  \\begin{figure*} \\centerline{\\includegraphics[width=9cm]{f1a.eps}\\includegraphics[width=9cm]{f1b.eps}} \\figcaption[simple]{The FIR-radio correlation at high redshift, in the rest-frame, using the B07 star formation law.  On left is the case with $B \\propto \\Sigma_g^{0.7}$ and no winds; on right, the case with $B \\propto \\rho^{0.6}$ and winds.  Solid lines have $h = 100~\\pc$ for starbursts ($\\Sigma_g \\ge 0.1~\\gcm2$; $\\Sigma_{\\rm SFR} \\ge 2 - 4~\\Msun~\\kpc^{-2}~\\yr^{-1}$), while dotted lines have $h = 1~\\kpc$ for starbursts.  Both panels show that a linear FRC is produced at $z = 0$ with the correct normalization.  Galaxies become radio-dim at high redshift as IC losses off the CMB increase.  Note that in each case the puffy starbursts do not lie on the same FRC as compact starbursts: there is a systematic offset caused by the unbalancing of the high-$\\Sigma_g$ conspiracy (\\S~\\ref{sec:Puffy}).  We do not include thermal radio emission, which will set a minimum radio luminosity or maximum $q^{\\prime}_{\\rm FIR}$ at each $\\Sigma_g$.\\label{fig:LFIRRadioRest}} \\end{figure*}  We show the rest-frame FRC for our model with $B \\propto \\Sigma_g^{0.7}$, the B07 star-formation law, and no winds, in Figure~\\ref{fig:LFIRRadioRest} (\\emph{left panel}) as an example.  The solid dark red line is the $z = 0$ FRC.  All solid lines assume that starbursts are compact, with $h = 100~\\pc$.  It is clear from Figure~\\ref{fig:LFIRRadioRest} that compact starbursts show little evolution in the FRC, while low surface density galaxies have lower radio luminosities at high redshift.  This behavior is robust in all of the variants.  The cause of evolution in the FRC is Inverse Compton losses off the CMB for CR electrons and positrons.  Figure~\\ref{fig:FRCEvolution} (\\emph{left panel}) shows that the FRC should display relatively little evolution out to $z \\approx 1$, except for the lowest surface brightness galaxies.  However, normal galaxies have suppressed synchrotron radio emission at $z \\approx 2$, a factor of $\\sim 2$ for $\\Sigma_g = 0.01~\\gcm2$ ($\\Sigma_{\\rm SFR} \\approx 0.06~\\Msun~\\kpc^{-2}~\\yr^{-1}$) and of order 10 for $\\Sigma_g = 0.001~\\gcm2$ ($\\Sigma_{\\rm SFR} \\approx 0.001 - 0.002~\\Msun~\\kpc^{-2}~\\yr^{-1}$).  The radio luminosities continue to fall with redshift.  At $z \\approx 5$, IC off the CMB starts to matter even for the weaker starbursts ($\\Sigma_g = 0.1~\\gcm2$; $\\Sigma_{\\rm SFR} \\approx 2 - 4~\\Msun~\\kpc^{-2}~\\yr^{-1}$).  Dense starbursts ($\\Sigma_g = 10~\\gcm2$; $\\Sigma_{\\rm SFR} \\approx 900~\\Msun~\\kpc^{-2}~\\yr^{-1}$ for the K98 law or $9000~\\Msun~\\kpc^{-2}~\\yr^{-1}$ for the B07 law) remain on a linear FRC even at $z \\approx 10$.  The strong cooling from bremsstrahlung, ionization, and IC off starlight implied by the high-$\\Sigma_g$ conspiracy (\\S~\\ref{sec:Theory}) acts as a buffer against IC losses off the CMB.  Similarly, diffusive escape provides a similar buffering effect in low-$\\Sigma_g$ galaxies, so that the increased IC losses off the CMB does not suppress the radio emission quickly.  In other words, most of the power from IC scattering of CMB photons does not come at the expense of synchrotron, but of other losses.  The radio dimming cannot be described by a simple competition between magnetic field energy density and the CMB energy density; the CMB energy density must also compete with \\emph{every other loss process}.  The buffering actually serves as an important test for both conspiracies described in \\S~\\ref{sec:Theory}.  In high-$\\Sigma_g$ galaxies and starbursts, we predict that bremsstrahlung, ionization, and Inverse Compton of starlight already take a large portion of a \\GHz~electron's energy budget; hence, the high-$\\Sigma_g$ conspiracy works to hold $L_{\\rm TIR}^{\\prime}/L_{\\rm radio}^{\\prime}$ constant out to quite high redshift.  The buffering essentially doubles the redshift that compact starbursts remain on the FRC; the weakest starbursts ($\\Sigma_g = 0.1~\\gcm2$; $\\Sigma_{\\rm SFR} \\approx 2 - 4~\\Msun~\\kpc^{-2}~\\yr^{-1}$) are radio-dim by a factor of $3$ at $z = 10$ with buffering, instead of $z = 5$ without the buffering.  In low-$\\Sigma_g$ galaxies, electrons and positrons can easily escape before they lose energy to Inverse Compton off the CMB at low enough $z$; thus, the low-$\\Sigma_g$ conspiracy also works to hold $L_{\\rm TIR}^{\\prime}/L_{\\rm radio}^{\\prime}$ constant.  The suppression due to IC losses off the CMB seen in Figure~\\ref{fig:LFIRRadioRest} can be estimated by examining the ratios of the synchrotron cooling time to the total loss time, including both escape and cooling losses.  In Appendix~\\ref{sec:CMBRadioDim}, we derive the ratios of loss times and we show that for a given Schmidt law, a critical redshift $z_{\\rm crit}$ can be defined for each $\\Sigma_g$ and $\\Sigma_{\\rm SFR}$ at which radio emission is suppressed.  We define $z_{\\rm crit}$ to be the redshift when the rest-frame radio luminosity at $\\nu^{\\prime} = 1.4~\\GHz$ is suppressed by a factor of 3 compared to $z = 0$.  For our model with the B07 star-formation law, no winds, and $B \\propto \\Sigma_g^{0.7}$, we find the critical redshifts are approximated as \\begin{equation} \\label{eqn:B07zCrit} z_{\\rm crit} \\approx \\left\\{ \\begin{array}{ll} 1.4 & ({\\rm Normal~galaxies}, \\Sigma_{\\rm SFR} \\la 0.02)\\\\ 5.8 \\Sigma_{\\rm SFR}^{0.23} - 1 & ({\\rm Normal~galaxies}, \\Sigma_{\\rm SFR} \\ga 0.02)\\\\ 5.7 \\Sigma_{\\rm SFR}^{0.23} - 1 & ({\\rm Puffy~starbursts})\\\\ 7.4 \\Sigma_{\\rm SFR}^{0.23} - 1 & ({\\rm Compact~starbursts}), \\end{array} \\right. \\end{equation} where $\\Sigma_{\\rm SFR}$ is in units of $\\Msun~\\kpc^{-2}~\\yr^{-1}$.  We refer the reader to Appendix~\\ref{sec:CMBRadioDim}, where we present similar relations for our other models of the FRC.  We note that the radio suppression could, conceivably, be used to measure the temperature of the CMB at high redshift.  In principle, this method could apply to any galaxy with a radio and FIR detection.  However, the conspiracies would have to be accounted for, and they are affected significantly by both the gas surface density $\\Sigma_g$ and scale height $h$ (see \\S~\\ref{sec:Puffy}).  Any measurement of the CMB temperature would depend on assumptions of the galaxy properties and would be model-dependent.  \\begin{figure*} \\centerline{\\includegraphics[width=9cm]{f2a.eps}\\includegraphics[width=9cm]{f2b.eps}} \\figcaption[simple]{The evolution of the FIR-radio correlation, both in the rest-frame at $\\nu^{\\prime} = 1.4~\\GHz$ (\\emph{left panel}), and as inferred for the rest-frame from observations at $\\nu = 1.4~\\GHz$ (\\emph{right panel}).  Black lines have $h = 100~\\pc$ for starbursts, while grey lines have $h = 1~\\kpc$ for starbursts.  Both panels show the case with $B \\propto \\Sigma_g^{0.7}$, no winds, and the B07 star-formation law.  Normal galaxies become radio dim because of IC losses off the CMB at intermediate redshift.  Both compact and puffy starbursts maintain their rest-frame radio luminosities until high redshift.  All galaxies are ``buffered'' by non-synchrotron losses (\\S~\\ref{sec:FRCEvolution} and Appendix~\\ref{sec:CMBRadioDim}), so that the evolution is not as great as would be expected with only synchrotron and IC losses off the CMB.  The inferred evolution is usually greater than the rest-frame evolution, because normal galaxies and puffy starbursts have $\\alpha \\ga 0.7$. \\label{fig:FRCEvolution}} \\end{figure*}  As Figure~\\ref{fig:FRCEvolution} (\\emph{right panel}) shows, the inferred rest-frame values of $L^{\\prime}_{\\rm TIR}/L^{\\prime}_{\\rm radio}$ show additional apparent evolution, simply because the spectral slopes of the galaxies are not exactly $0.7$.  Compact starbursts have flatter spectra\\footnote{In this paper (unlike LTQ), $\\alpha^{\\prime} \\equiv \\frac{d \\log F^{\\prime}_{\\nu}}{d \\log \\nu^{\\prime}}$ and $\\alpha \\equiv \\frac{d \\log F_{\\nu}}{d \\log \\nu}$ are instantaneous spectral slopes, not the measured spectral slopes between two observed frequencies, unless otherwise noted. \\label{ftnt:AlphaNotation}} with $\\alpha \\approx 0.4 - 0.6$ at 1.4 GHz, so by adopting $\\alpha = 0.7$ they \\emph{appear} to become slightly radio brighter until $z \\approx 1$, after which at higher frequencies their spectra steepen, and their radio emission appears to dim again at higher $z$.  Normal galaxies have steeper spectra with $\\alpha \\approx 0.9 - 1.0$ at 1.4 GHz, so their apparent radio luminosity sinks below their true radio luminosity at high redshift.  At higher redshifts, our models predict that galaxies have intrinsically steeper radio spectra, because of increased IC losses from the CMB.  The rest-frame spectral slope $\\alpha^{\\prime}$ of normal galaxies at 1.4 GHz and asymptotes at $\\sim 1.1$ by $z \\approx 2 - 3$, when losses are dominated by IC for our injection spectrum of CRs with $p = 2.2$.  There is much less intrinsic rest-frame evolution of starburst radio spectra; even at $z \\approx 5$, $\\alpha^{\\prime}$ increases only by 0.1 for the weakest compact starbursts.  The observable $\\alpha^{1.4~\\GHz}_{610~\\MHz}$ shows much more pronounced evolution with redshift for starbursts, increasing by $0.1$ to $z \\approx 1 - 2$, and $0.2$ to $z \\approx 4 - 5 $.  This effect arises simply because we predict the CR electron and positron spectra steepen with rest-frame frequency as synchrotron and IC losses become stronger \\citep[e.g.,][LTQ]{Thompson06}: at higher redshift and fixed observing frequency, we are seeing higher energy electrons and positrons.  An important effect that we do not consider is the thermal free-free radio emission.  This will set a minimum total observed radio emission that is directly proportional to the star-formation luminosity.  Thus, the true total radio deficit at GHz will not be as big as the synchrotron-only deficits plotted in Figures~\\ref{fig:LFIRRadioRest} and \\ref{fig:FRCEvolution}.  Thermal free-free emission will also flatten the spectrum, especially at high frequencies.  However, free-free emission is much fainter than the synchrotron luminosity at GHz frequencies, except in the faintest star-forming galaxies \\citep{Condon92,Hughes06}.  \\subsection{The Radio Excess (or Deficit) of Puffy Starbursts: Submillimeter Galaxies} \\label{sec:Submillimeter} \\label{sec:Puffy} As shown in Figure~\\ref{fig:LFIRRadioRest}, puffy starbursts fall on a linear FIR-radio correlation of their own (dotted lines), in line with observations of SMGs \\citep{Kovacs06,Murphy09}.  The radio luminosity of these galaxies is nonetheless the result of a conspiracy, between IC losses on starlight, which decrease the radio luminosity, and the enhanced radio emission from secondary electrons and positrons, and the $\\nu_C$ effect (\\S~\\ref{sec:Theory}).  The variation in $L_{\\rm TIR}^{\\prime}/L_{\\rm radio}^{\\prime}$ for puffy starbursts alone is usually less than a factor of $2$ over the range $0.1~\\gcm2 \\le \\Sigma_g \\le 10~\\gcm2$ (see Table~\\ref{table:Models}; in most variants the variation is $\\sim 1.6$).  Like compact starbursts, escape plays essentially no role in most of the models, except that winds can slightly decrease the radio emission in relatively tenuous $\\Sigma_g = 0.1~\\gcm2$ ($\\Sigma_{\\rm SFR} \\approx 2 - 4~\\Msun~\\kpc^{-2}~\\yr^{-1}$) starbursts (footnote~\\ref{ftnt:WindVariant}).   \\subsubsection{$B \\propto \\Sigma_g^{0.7-0.8}$: Radio-Bright Puffy Starbursts} It is plain from Figure~\\ref{fig:LFIRRadioRest} (\\emph{left panel}) that the normalization of the puffy starburst FRC (dotted lines) is different than the FRC of the compact starbursts and normal galaxies (solid lines) when $B \\propto \\Sigma_g^{0.7}$.  We show in Table~\\ref{table:Models} that models with $B \\propto \\Sigma_g^{0.7-0.8}$ have radio-bright puffy starbursts compared to the observed local FRC, by a factor of $2 - 4$ at $z = 0$.  Like the compact starbursts, puffy starbursts show little rest-frame evolution in $L_{\\rm TIR}^{\\prime}/L_{\\rm radio}^{\\prime}$, except at relatively low surface densities ($\\Sigma_g = 0.1~\\gcm2$; $\\Sigma_{\\rm SFR} \\approx 2 - 4~\\Msun~\\kpc^{-2}~\\yr^{-1}$), where they become radio dim at high redshifts because of IC losses on CMB photons.  Therefore, we predict that puffy starbursts, which are mainly observed at high $z$, have \\emph{intrinsically} different radio properties not caused by their redshift.  We propose a natural explanation of this radio excess in the framework of LTQ (\\S~\\ref{sec:Theory}).  In dense starbursts, protons are efficiently converted into secondary electrons and positrons through inelastic proton-proton scattering, which contribute to the synchrotron emission.  Furthermore, the $\\nu_C$ effect increases $L_{\\rm TIR}^{\\prime}/L_{\\rm radio}^{\\prime}$ for starbursts which have larger $B$.  Compact starbursts lying on the $z = 0$ FRC balance these effects with increased bremsstrahlung, ionization, and IC losses, which compete with the synchrotron losses and suppress the excess radio luminosity (\\S~\\ref{sec:Theory}).  In puffy starbursts with relatively low volume densities compared to their compact cousins at fixed $\\Sigma_g$, however, bremsstrahlung and ionization are not strong enough to compensate for these effects.  Only synchrotron and IC losses remain, upsetting the conspiracy.  The radio excess predicted by this picture is systematically greater when using the K98 star-formation relation, because the IC loss rate on starlight is smaller at fixed surface density by a factor of 2.5 to 11 from weak starbursts to the densest starbursts.  This freedom to vary the radio-excess with $B$ does not exist in the standard calorimeter model, where all CR electron/positron energy goes into synchrotron emission, and the radio emission saturates.  The balance between synchrotron and the other forms of cooling can be changed in starbursts to alter the normalization of the FRC, even though escape is negligible.  There is a reason to expect our allowed $B \\propto \\Sigma_g^{0.7-0.8}$ specifically: radiation pressure may drive turbulence and enhance the magnetic field until its energy density is comparable to radiation \\citep[e.g.,][]{Thompson08}.  If the K98 relation holds, then since the magnetic energy density scales as $U_B \\propto U_{\\rm ph}$, $B \\propto \\Sigma_g^{0.7}$, and if the B08 relation holds, then $B \\propto \\Sigma_g^{0.85}$.  This explanation is more problematic for the B07 relation and $B \\propto \\Sigma_g^{0.7}$. \\footnote{However, we have scaled $B$ to the \\emph{total} Milky Way magnetic field strength of $6~\\muGauss$ near the Solar Circle.  If most of the magnetic field strength in starbursts is driven by turbulence, perhaps scaling to the disordered Galactic magnetic field strength \\citep[$\\sim 2 \\muGauss$;][]{Beck01} near the Solar Circle would make more sense, because the ordered magnetic field arises from a different process.  The lower normalization for $B$ then would partly compensate for the steeper dependence on $\\Sigma_g$.}  However, while we assume that there is a parametrization for $B$ that applied to both compact starbursts and puffy starbursts, the radio excess should arise more generally.  The radio excess arises simply because synchrotron cooling time is shorter than the bremsstrahlung and ionization cooling times in puffy starbursts, but longer in compact starbursts at fixed $\\Sigma_g$: at fixed $\\nu^{\\prime}$, $t_{\\rm brems} \\propto n^{-1} \\propto h / \\Sigma_g$ and $t_{\\rm ion} \\propto B^{-1/2} n^{-1} \\propto h / (B^{1/2} \\Sigma_g)$.  We therefore expect there to be a radio excess with respect to the FRC in any puffy starburst with a strong enough magnetic field, with the exact enhancement depending on magnetic field strength because of the remaining competition, after ionization and bremsstrahlung are sub-dominant, from the IC losses on starlight.\\footnote{If the magnetic field strength does not go very roughly as $B \\propto \\Sigma_g^{0.7}$, however, there will be more scatter in $L^{\\prime}_{\\rm TIR}/L^{\\prime}_{\\rm radio}$ for puffy starbursts at fixed $h$ and $z$.  SMGs would not form their own FRC if $B$ was the same for all SMGs regardless of $\\Sigma_g$, or if $B$ increased \\emph{very} steeply with $\\Sigma_g$, because there still must be a conspiracy with IC losses off starlight.  Since SMGs do appear to form their own FRC \\citep[e.g.,][]{Murphy09,Michalowski09}, this could be evidence specifically that their magnetic field strengths increase with $\\Sigma_g$ (or $\\rho$).}  Our results imply that a moderate radio excess at the factor of $\\sim 3$ level alone is not a safe indicator of the presence of a radio-loud AGN, especially at high redshifts where SMGs are observed.  While radio excess with respect to the local FRC has been suggested as a selection criterion for radio-loud AGNs \\citep[e.g.,][]{Yun01,Yang07,Sajina08}, our models with $\\Sigma_g^{0.7-0.8}$ imply that $q_{\\rm FIR} \\approx 1.7 - 2.0$ for puffy starbursts powered by star-formation alone.  A radio excess is inexplicable in our models only when the source is an order of magnitude brighter ($q_{\\rm FIR} \\la 1.5$) in the radio than predicted from the FIR emission.  While SMGs are relatively rare and may not be a problem in small samples, we recommend that other means be used to be sure that the radio-excess is caused by an AGN, such as a flat radio spectrum, radio morphology, mid-IR colors, or the presence of strong X-ray emission \\citep[see also][]{Murphy09}.  \\subsubsection{$B \\propto \\rho^{0.5 - 0.6}$: Radio-Dim Puffy Starbursts} A magnetic field dependence of $B \\propto \\rho^{0.5}$ appears to hold for Galactic molecular clouds \\citep[e.g.,][]{Crutcher99}, and LTQ found that the FIR-radio correlation was consistent with this magnetic field dependence.  The existence of galaxies with different scale heights allows us to distinguish the two possibilities for magnetic field scaling.  If $B \\propto \\Sigma_g^a$, then the magnetic field strength will be the same for all galaxies with the same $\\Sigma_g$, regardless of scale height.  In models with $B \\propto \\rho^{0.5 - 0.6}$, by contrast, the magnetic field strength is weaker in puffy starbursts than in compact starbursts with the same $\\Sigma_g$.  As seen in Figure~\\ref{fig:LFIRRadioRest} (\\emph{right panel}, dotted lines), puffy starbursts again form their own FRC.  In models with $B \\propto \\rho^{0.5 - 0.6}$, they are radio dim compared to the $z \\approx 0$ FRC.  We show in Table~\\ref{table:Models} that the normalization of the FRC is radio-dim by a factor of $\\sim 1.2 - 2.0$.  We can explain this in the LTQ theory of the FRC as well.  The magnetic field strength must increase more slowly with $\\rho$ than with $\\Sigma_g$ to reproduce the $z \\approx 0$ FRC, because $h \\propto \\Sigma_g / \\rho$ is 10 times smaller in compact starbursts than normal galaxies and puffy starbursts.  Compact starbursts are highly compressed with respect to normal galaxies, so they have strong magnetic fields and synchrotron radio emission is strong enough to compete with the other losses.  Puffy starbursts are not compressed, so that their magnetic fields are weak and synchrotron losses cannot keep up with IC losses, nor with bremsstrahlung and ionization as 1.4 GHz emission traces ever lower electron energies at higher magnetic field strengths.  Puffy starbursts therefore turn out to be radio dim compared to compact starbursts on the $z \\approx 0$ FRC, if $B \\propto \\rho^{0.5 - 0.6}$.  As before, the B07 star-formation relation predicts greater IC losses and therefore weaker radio emission.  Because of the claims in the literature that SMGs are radio bright, we do not favor these models.  The suggested relative radio brightness of high-$z$ SMGs therefore provides some evidence that, in fact, $B \\propto \\Sigma_g^{0.7 - 0.8}$ rather than $B \\propto \\rho^{0.5 - 0.6}$.  However, the matter of whether high-$z$ SMGs are in fact radio bright is not yet settled.  Although LTQ concluded that $B$ must increase dramatically from normal galaxies to dense starbursts (\\S~\\ref{sec:Theory}), they were unable to distinguish between these two possibilities with the $z \\approx 0$ FRC alone.  For this reason, high-$z$ starbursts and their qualitatively different morphologies compared to those at $z \\approx 0$ can distinguish theories of the FRC.  \\subsubsection{Spectral slopes} A prediction of all of our variants is that puffy starbursts like submillimeter galaxies should have steep non-thermal radio spectra, with $\\alpha \\approx 0.8 - 1.0$ (see Table~\\ref{table:Models}).  The steep spectra are caused by strong synchrotron cooling in the $B \\propto \\Sigma_g^{0.7-0.8}$ case and the relatively stronger IC cooling off starlight in the $B \\propto \\rho^{0.5 - 0.6}$ case.  In general, puffy starbursts should have roughly the same $\\alpha$ as normal galaxies in the local universe, which tends to be somewhat higher ($\\alpha \\approx 0.7 - 1.0$) than in compact starbursts ($\\alpha \\la 0.7$).  The slope should hold even out to extremely high $\\Sigma_g$, as long as starbursts are puffy.  In contrast, we find that $\\alpha \\approx 0.5$ in compact starbursts, because of efficient ionization and bremsstrahlung losses, which flatten the equilibrium CR spectrum because of their energy dependence.  As we note in LTQ, our predicted spectral index for normal galaxies is somewhat too high, and this difference in $\\alpha$ may carry over to the puffy starbursts.  However, the significant difference in $\\alpha$ between compact and puffy starbursts should remain as a general prediction of our model: compact starbursts should have flatter spectra than puffy starbursts.  The high spectral slopes can be observed either with direct measurements of multifrequency data of individual submillimeter galaxies, or with single frequency observations at a variety of redshifts.  There are relatively few measurements of $\\alpha$ for submillimeter galaxies specifically; faint radio sources have $\\alpha \\approx 0.5 - 0.7$ \\citep{Huynh07,Bondi07}, though that sample includes both compact starbursts and AGNs.  \\citet{Sajina08} do find that $\\alpha_{610~\\MHz}^{1.4~\\GHz} \\approx 0.8$ for SMGs, comparable to our predictions.  They also find that submillimeter galaxies have a radio-excess, in agreement with Figure~\\ref{fig:LFIRRadioRest}.  More recently, \\citet{Ibar10} found an average $\\alpha_{610~\\MHz}^{1.4~\\GHz} \\approx 0.75 \\pm 0.06$, which is somewhat flatter than our models.  These spectral slopes are not different from normal star-forming galaxies, but are noticeably steeper than local ULIRGs \\citep{Clemens08}.  However, we do not account for free-free absorption, which probably flattens the spectra of local ULIRGs like Arp 220 at low frequency \\citep{Condon91}, and is not well understood in SMGs.  Since puffy starbursts have steeper spectra than compact starbursts, we expect their inferred $L_{\\rm TIR}^{\\prime}/L_{\\rm radio}^{\\prime}$ will increase with redshift: if the true radio spectral slopes of SMGs are greater than the assumed $\\alpha$ by $\\Delta\\alpha$, they will appear to become radio dimmer by a factor $(1 + z)^{\\Delta \\alpha}$, or up to $\\sim 40\\%$ at $z = 2$.  \\begin{figure} \\centerline{\\includegraphics[width=9cm]{f3.eps}} \\figcaption[figure]{The rest-frame synchrotron radio spectra of starbursts with $\\Sigma_g = 1~\\gcm2$, using the B07 star-formation law. Puffy starbursts are dashed, while compact starbursts are solid. These spectra do not include thermal absorption (at low frequencies, $\\la 1\\ \\GHz$) or emission (at high frequencies, $\\ga 30\\ \\GHz$).  \\label{fig:StarburstSpectra}} \\end{figure}  In Figure~\\ref{fig:StarburstSpectra}, we show the expected radio synchrotron spectra of starburst galaxies, without correcting for thermal absorption or thermal emission.  At a rest-frame frequency of 1 GHz, puffy starbursts (dashed) have steeper radio spectra than compact starbursts (solid).  Note that at high frequencies ($\\nu^{\\prime} \\ga 10\\ \\GHz$), the ratio of the radio luminosities per unit star formation of the compact and puffy starbursts asymptotes to a value set by the ratio of $U_B$ and $U_{\\rm ph}$ in these starbursts.  At these high frequencies, only synchrotron and IC cooling are effective, and IC cooling would be the same for puffy and compact starbursts because of the Schmidt Law (\\S~\\ref{sec:Theory}).  For $B \\propto \\Sigma_g^a$, $U_B$ is the same for puffy and compact starbursts, but for $B \\propto \\rho^a$, puffy starbursts have much smaller $U_B$.  Thus, measurements of the synchrotron radio emission of SMGs at high $\\nu^{\\prime}$ could determine the magnetic field strength of SMGs and determine which scenario applies.  Of course, there are unlikely to be two perfectly distinct populations of compact starbursts and puffy starbursts.  Instead, there may be a continuum variation in scale heights from tens to thousands of parsecs.  We would then expect to see a larger scatter, both in $L_{\\rm TIR}^{\\prime}/L_{\\rm radio}^{\\prime}$ and $\\alpha$, in a full sample of both the most compact and the most puffy starbursts.  \\citet{Murphy09} find that submillimeter galaxies do have a larger scatter in $q_{\\rm TIR}^{\\prime}$ than other galaxies.  However, \\citet{Ibar10} find a relatively small scatter of $\\sim 0.3$ in SMG radio spectral index.  Importantly, for larger $h$, both $L_{\\rm TIR}^{\\prime}/L_{\\rm radio}^{\\prime}$ and $\\alpha^{\\prime}$ asymptote as CR electron and positron losses are entirely determined by synchrotron and IC; $h \\approx 1~\\kpc$ starbursts are already near this limit.  In other words, for arbitrarily large $h$, the radio excess with respect to the $z \\approx 0$ FRC asymptotes to a value of $\\sim 5 - 10$, depending on the assumed Schmidt Law, and the radio spectral slope asymptotes to $\\sim 1.1$ for $p = 2.2$.  \\subsection{Synchrotron Radio Emission as a Star-Formation Tracer} \\label{sec:SFTracer}  \\begin{figure} \\centerline{\\includegraphics[width=9cm]{f4.eps}} \\figcaption[figure]{The relationship between the inferred rest-frame radio emission and star-formation, using the B07 star formation law, $B \\propto \\Sigma_{\\rm SFR}^{0.7}$, and no winds.  Solid lines have $h = 100~\\pc$ for starbursts, while dotted lines have $h = 1~\\kpc$ for starbursts.  A flat line would indicate that radio emission is directly proportional to $\\Sigma_{\\rm SFR}$.  We see that at low $\\Sigma_g$, the radio flux underestimates the star-formation rate even at $z = 0$, because of CR electron escape (the low-$\\Sigma_g$ conspiracy of \\S~\\ref{sec:Theory}).  The radio flux overestimates the star-formation rate for puffy starbursts, because the high-$\\Sigma_g$ conspiracy is unbalanced (\\S~\\ref{sec:Theory}; \\S~\\ref{sec:Puffy}).  Finally, the radio emission is suppressed at high redshift, partly because of IC losses off the CMB, and partly because $\\alpha \\ga 0.7$ for puffy starbursts and normal galaxies ($\\alpha = 0.7$ assumed here for the k-correction).\\label{fig:SFRRadio}} \\end{figure}  Relying on the FRC, a number of studies have used the GHz radio emission as a tracer of star-formation \\citep[e.g.,][]{Cram98,Mobasher99,Haarsma00,Carilli08,Seymour08,Garn09}.  Radio emission has the advantage that it is unaffected by dust obscuration, making it potentially very useful in starbursts, and therefore for most of the star-formation at $z \\ga 1$ \\citep[e.g.,][]{Chary01,LeFloch05,Dole06,Magnelli09,Pascale09}.  Our models let us evaluate the theoretical basis for radio as a SFR indicator at all redshifts.  We show the predicted radio emissivity as a function of star formation in Figure~\\ref{fig:SFRRadio}.  At low surface densities ($\\Sigma_g \\la 0.01~\\gcm2$; $\\Sigma_{\\rm SFR} \\la 0.06~\\Msun~\\kpc^{-2}~\\yr^{-1}$), synchrotron radio has a non-linear dependence on star-formation at all redshifts.  The weak radio emission is caused by electrons escaping their host galaxies, as inferred by \\citet{Bell03} and discussed in LTQ.  That is, normal galaxies are not perfect electron calorimeters, so the radio emission is not a reliable star-formation tracer at low $\\Sigma_{\\rm SFR}$.  At higher redshift, synchrotron emission is diminished by Inverse Compton off the CMB.  For a Milky Way-like galaxy, we find that radio is a good star-formation tracer at $z \\la 1$, but underestimates it significantly by $z \\ga 2$ (eq. \\ref{eqn:B07zCrit}).  Already galaxies with star formation rates similar to Galactic levels are beginning to be observed in the radio at high redshift \\citep{Garn09}, so that the IC suppression may soon be observed.  However, IC losses off the CMB should not be important even in the weakest starbursts until $z \\ga 4$, as is also visible by the redshift evolution of the FRC in Figure~\\ref{fig:FRCEvolution}.  Radio emission does grow linearly with star-formation rate between normal galaxies with $\\Sigma_g = 0.01~\\gcm2$ (at $z \\la 2$) and compact starbursts.  Therefore, it serves as an acceptable star-formation indicator for these galaxies.  However, if $B \\propto \\Sigma_g^{0.7-0.8}$, puffy starbursts like SMGs have about $2 - 4$ times the radio emission at any given star-formation rate than compact starbursts.  Therefore, we expect that if $B \\propto \\Sigma_g^{0.7-0.8}$, the usual radio emission estimate based on the $z \\approx 0$ FRC will \\emph{overestimate} their star formation rates by a factor of $\\sim 2 - 4$.  The excess is greatest at $\\Sigma_g \\approx 1~\\gcm2$ corresponding to $\\Sigma_{\\rm SFR} \\approx 40 - 200~\\Msun \\kpc^{-2} \\yr^{-1}$, typical of observed SMGs.  Among the puffy starbursts themselves, the radio emission grows linearly with star-formation rate.  If instead $B \\propto \\rho^{0.5 - 0.6}$, the radio emission will underestimate the star-formation rate.  However, because the SMGs lie on their own FRC, there is little real redshift evolution in the radio emissivity of puffy starbursts, because of the buffering provided by IC losses off starlight (see \\S~\\ref{sec:FRCEvolution}, Appendix~\\ref{sec:CMBRadioDim}).  Assuming a spectral slope $\\alpha = 0.7$ will also underestimate the radio emissivity, since these galaxies can have steep radio spectra.  This explains the apparent evolution with $z$ of puffy starbursts in Figure~\\ref{fig:SFRRadio}: we have applied a k-correction using a typically employed $\\alpha = 0.7$ when, in fact, the true synchrotron spectra are steeper.  If the radio star-formation tracer could be calibrated to the special conditions in puffy starbursts like SMGs, taking into account their different scale height, radio emissivity, and spectral slopes, we predict that radio would be a more accurate star-formation tracer for them.  \\begin{figure} \\centerline{\\includegraphics[width=9cm]{f5.eps}} \\figcaption[figure]{The \\emph{observed} 1.4 GHz synchrotron radio flux per unit star-formation as a function of redshift, using the B07 star formation law, $B \\propto \\Sigma_{\\rm SFR}^{0.7}$, and no winds.  EVLA will have a sensitivity of $\\sim 1\\ \\muJy$ and SKA will have a sensitivity of $\\sim 20\\ \\nJy$.\\label{fig:GHzRadioFluxes}} \\end{figure}  We can also directly calculate the \\emph{observed} GHz radio flux density $S_{\\nu}$ from synchrotron emission of star-forming galaxies.  We show the predicted flux density per unit star-formation at observer-frame 1.4 GHz in Figure~\\ref{fig:GHzRadioFluxes}.  Surveys with the Expanded Very Large Array (EVLA) will have a continuum sensitivity of $\\sim 1\\ \\muJy$ at frequencies of 1 - 50 GHz, and it should be able to directly detect galaxies with $\\Sigma_g \\ga 0.01\\ \\gcm2$ and star-formation rates of $1\\ \\Msun\\ \\yr^{-1}$ out past $z \\approx 0.5$.  As stated in \\citet{Murphy09c}, a starburst like M82 with SFR $\\approx 3\\ \\Msun\\ \\yr^{-1}$ will become undetectable past $z \\approx 1$.  However, the buffering effect we emphasize in this paper preserves the radio emission of dense starbursts at high $z$, so that bright starbursts will be detectable further: starbursts with SFR $\\ga 100\\ \\Msun\\ \\yr^{-1}$ will be detectable with EVLA at 1.4 GHz out to $z \\approx 4 - 5$ and the most intense starbursts (SFR $\\ga 1000\\ \\Msun\\ \\yr^{-1}$ and $\\Sigma_g \\ga 1~\\gcm2$) will be detectable out past $z \\approx 10$ in synchrotron emission.  \\citet{Murphy09c} predicts the Square Kilometer Array (SKA) will be sensitive to star-forming galaxies with flux densities of $\\sim 20\\ \\nJy$.  If this sensitivity is attained, then Milky Way-like galaxies ($\\Sigma_g \\approx 0.01\\ \\gcm2$; SFR $\\approx 1\\ \\Msun\\ \\yr^{-1}$) will be directly detectable at 1.4 GHz in synchrotron emission out to $z \\approx 2$, and starbursts with SFR $\\approx 1\\ \\Msun\\ \\yr^{-1}$ will be detectable at 1 GHz beyond $z \\approx 3$.  Even at $z \\approx 10$, the synchrotron emission of dense, compact starbursts ($\\Sigma_g \\ga 1\\ \\gcm2$) with SFR greater than $10\\ \\Msun\\ \\yr^{-1}$ should be detectable at 1.4 GHz with SKA (models with $B \\propto \\rho^{0.5-0.6}$ have radio-dim puffy starbursts, and these are detectable at $z = 10$ with the SKA only for SFR greater than $20 - 60\\ \\Msun\\ \\yr^{-1}$).  By contrast, \\citet{Murphy09c} found a sensitivity of $25\\ \\Msun\\ \\yr^{-1}$, based on the free-free emission; this limit will apply to normal galaxies and weak starbursts where the synchrotron emission is suppressed.  These sensitivities assume that natural confusion, in which radio sources overlap, will not hamper the SKA; estimates for the natural confusion limit vary from nJy to $\\muJy$ levels \\citep[e.g.,][]{Jackson04,Condon09,Murphy09c}.  SKA and EVLA will also have good spectral coverage, which may help measurements of the spectral index.  If a galaxy is detected at $5 \\sigma$ ($\\sim \\muJy$ for EVLA and $20\\ \\nJy$ for SKA; \\citealt{Murphy09c}) at two different frequencies $\\nu_1$ and $\\nu_2$ with $\\nu_2 / \\nu_1 = 5$, then $\\alpha_1^2$ can be constrained to $\\sim 0.1 - 0.15$ at the $1 \\sigma$ level.  EVLA will be better at high observer-frame frequencies (1 - 50 GHz).  A problem for the EVLA will be the increasing fraction of thermal emission, which is expected to dominate the emission of starbursts at $\\nu^{\\prime} \\approx 30\\ \\GHz$.  EVLA will therefore not easily measure the nonthermal spectral indices of galaxies at high z.  On the other hand, the SKA will face free-free absorption when observing low-z starbursts; for example, Arp 220 may be optically thick even at 1 GHz \\citep{Condon91}.  At high redshift, however, the rest-frame frequencies SKA will observe will be less affected by free-free absorption.  \\subsection{SMGs, The Radio Background, and Radio Source Counts} \\label{sec:RadioBackground} Star-formation in galaxies over cosmic time produces a diffuse radio synchrotron background.  Our models indicate that submillimeter galaxies and other puffy starbursts ought to have enhanced radio emission.  In principle, this means that the synchrotron radio background could be up to $\\sim 2 - 4$ times higher than usually predicted from the Cosmic Infrared Background (CIB) and a naive application of the $z \\approx 0$ FRC.  The magnitude of this implied radio excess is interesting, because ARCADE2 recently reported an excess radio background at 3~\\GHz, about five times higher than expected from star formation \\citep{Fixsen09,Seiffert09}.  \\citet{Singal09} found that constraints on Inverse Compton emission require the excess to come from regions with galactic-level ($\\ga \\muGauss$) magnetic fields, and suggest an evolution of the FRC as the source of the reported excess.  However, while SMGs are individually very bright and contribute much to the cosmic star-formation rate at $z \\ga 2.5$, they are not typical starbursts \\citep{Bavouzet08}.  Instead, they seem to represent a transient phase that can survive about 100~\\Myr~before their gas is depleted \\citep{Tacconi06,Pope08}.  We can estimate the total radio background enhancement by scaling to the contribution of SMGs to the CIB, which is largely reprocessed  starlight from galaxies at $z \\ga 1$ \\citep[e.g.,][]{Dole06,Devlin09}, and adjusting by the SMG radio excess.  Submillimeter galaxies do provide the majority of light at $\\sim 850 \\mum$, but this is only a small fraction of the total IR background.  At the peak of the CIB ($\\sim 160 \\mum$), submillimeter galaxies provide $\\la 10\\%$ of the total power, and possibly only $\\sim 2\\%$ \\citep{Chapman05,Dye07}.  Optimistically, the radio excess from SMGs would be $10\\% \\times (4 - 1) \\approx 30\\%$.  This is significant, but not enough to explain the very large ARCADE2 excess.  More conservatively, the excess is more likely $5\\% \\times 2 \\approx 10\\%$, and could be as little as a few percent.  Since the number of radio sources down to several $\\muJy$ is well known, we can estimate the fraction of the expected radio background comes from bright SMGs.  \\citet{Dole06} find a TIR background of about $24~\\nWBackUnits$; from the normalization of the FIR-radio correlation we expect a 1.4 GHz background of $\\nu I_{\\nu} \\approx 2.6 \\times 10^{-5}~\\nWBackUnits$.  \\citet{Chapman05} found an average radio flux density of $S_{1.4} \\approx 75\\ \\muJy$ for bright SMGs ($S_{\\rm 850\\ \\mu m} \\ga 5\\ \\mJy$).  The number counts of $75\\ \\muJy$ sources imply that they have a density of $\\sim 1500\\ \\deg^{-2}$, contributing roughly $\\sim 5.1 \\times 10^{-6}~\\nWBackUnits$ to the 1.4 GHz radio background \\citep[e.g.,][]{Gervasi08}.  Bright SMGs have an approximate density of $\\sim 600\\ \\deg^{-2}$ \\citep[e.g.,][]{Wang04,Coppin06}, so they constitute about $\\sim 40\\%$ of the background from $75\\ \\muJy$ sources, or $\\sim 8\\%$ of the expected 1.4 GHz background from star-formation.  This is roughly in line with our estimate of $\\sim 10\\%$, although the uncertainties are large enough that it could be consistent with SMGs lying on the FRC.  In any case, it is fairly clear that bright SMGs are not the source of the ARCADE excess.  The total excess could be greater if most starbursts at $z \\ga 1$ are puffy, and not just SMGs.  We do not expect this simply because most current studies show that the local FRC \\emph{does} hold out to high redshift for most observed galaxies and starbursts \\citep[e.g.,][]{Murphy09,Younger09,Garn09}.  The spectral slope would also be a problem: ARCADE2 inferred $\\alpha = 0.6$, while we predict a spectral slope $\\alpha \\ga 0.7$, steeper than local compact starbursts.  Resolved radio sources in the range $50\\ \\muJy \\la S_{1.4} \\la 1\\ \\mJy$, which are expected to be star-forming galaxies \\citep{Danese87,Condon89,Benn93}, contribute about $1.1 \\times 10^{-5}\\ \\nWBackUnits$ to the 1.4 GHz radio background \\citep{Gervasi08}, which is consistent with the total energetics expected from the FRC.  \\citet{deZotti10} also find consistency between the FRC and the observed radio source counts at $\\sim 30\\ \\muJy$.  Stacking studies of fainter radio sources have given somewhat ambiguous results on whether the FRC applies \\citep{Boyle07,Beswick08}, but \\citet{Garn09b} find no large evolution in the FRC down to $S_{1.4} \\approx 20\\ \\muJy$.  Any large radio excess would have to come either from a new population of low luminosity galaxies or very high-$z$ galaxies (see Figure~\\ref{fig:GHzRadioFluxes}); extrapolations of the higher flux source populations do not predict a large radio excess \\citep[e.g.][]{Gervasi08}.  Nonetheless, our work indicates that a radio excess from SMGs can be significant.  Conversely, if the luminosity function of galaxies was very steep, most galaxies could be intrinsically radio-dim with respect to the FRC, because of IC losses off the CMB (see \\S~\\ref{sec:FRCEvolution} and Figures~\\ref{fig:LFIRRadioRest} and~\\ref{fig:FRCEvolution}).  Then the FRC would overestimate the strength of the radio background.  However, most star-formation at high $z$ is believed to have occurred in starbursts \\citep{LeFloch05,Dole06}, so this possibility is unlikely.   We have applied the theory of LTQ to predict the FRC for redshifts $0 \\le z \\le 10$.  We use one-zone models of galaxies and starbursts with CR injection, cooling, and escape to predict the equilibrium, steady-state radio spectra of galaxies and starbursts over the entire range of the FRC.  Our goals were to determine how and why the low- and high-$\\Sigma_g$ conspiracies crucial to the $z \\approx 0$ FRC (\\S~\\ref{sec:Theory}) affect the FRC at high redshift, and to provide a quantitative model for predicting the critical redshifts at which galaxies deviate from the $z \\approx 0$ FRC.  We find the following:  \\begin{enumerate} \\item For compact starbursts ($h \\approx 100~\\pc$), we find relatively little evolution in the FIR-radio correlation out to $z \\approx 5 - 10$ (Figure~\\ref{fig:LFIRRadioRest}).  This is partly because the magnetic energy density in galaxies is strong enough to dominate the CMB even at high redshifts.  However, the high-$\\Sigma_g$ conspiracy (\\S~\\ref{sec:Theory}) also acts as a buffer against IC losses off the CMB; the increased IC losses must compete with the already present bremsstrahlung, ionization, and IC off starlight in addition to synchrotron losses.  The rest-frame radio spectral slope $\\alpha^{\\prime}$ at fixed $\\nu^{\\prime}$ does not change with $z$, but the observed $\\alpha$ at fixed $\\nu$ increases because the non-thermal synchrotron radio spectrum steepens at higher rest-frame frequency.  \\item We derive in Appendix~\\ref{sec:CMBRadioDim} the critical redshifts when Inverse Compton losses off the CMB suppress the radio luminosity of galaxies compared to the $z \\approx 0$ FRC.  These relations are given for our standard model in equation \\ref{eqn:B07zCrit}.  The non-thermal radio luminosity is suppressed severely in Milky Way-like galaxies ($\\Sigma_{\\rm SFR} \\approx 0.06~\\Msun~\\kpc^{-2}~\\yr^{-1}$) at $z \\approx 2$ and the weakest compact starbursts at $z \\ga 5$.  The spectrum at GHz steepens to $\\alpha \\approx 1$ because of these enhanced IC losses.  Nonetheless, the low-$\\Sigma_g$ conspiracy (\\S~\\ref{sec:Theory}) also acts to prevent the radio emission from steeply falling with redshift, since Inverse Compton losses off the CMB must be more efficient than diffusive escape, not just synchrotron losses (see \\S~\\ref{sec:FRCEvolution}).  \\item LTQ found that the $z = 0$ FRC demands that $B$ scales with $\\rho$ or $\\Sigma_g$ in galaxies lying on the Schmidt law.  In models with $B \\propto \\Sigma_g^{0.7-0.8}$, we find that puffy starbursts with $h = 1~\\kpc$ such as SMGs are radio bright compared to the $z = 0$ FRC by a factor of $\\sim 2 - 4$.  This follows from a breakdown of the high-$\\Sigma_g$ conspiracy (\\S~\\ref{sec:Theory}): bremsstrahlung and ionization cooling are weak in puffy starbursts relative to the compact starbursts that predominate in the $z = 0$ universe.  In contrast, in models with $B \\propto \\rho^{0.5 - 0.6}$, we find that puffy starbursts are radio dim compared to the observed FRC, because of weak synchrotron cooling relative to the IC losses.  Since several studies have reported radio excesses for SMGs, we favor the $B \\propto \\Sigma_g^{0.7-0.8}$ scaling; however the issue of whether SMGs are radio-bright is still not fully resolved.  In either case, puffy starbursts show little true evolution with $z$, though they may appear to have fainter rest-frame radio luminosities at high $z$ because of their steep spectra.  Puffy starbursts inevitably have high $\\alpha$ ($\\ga 0.7$), since bremsstrahlung and ionization losses are weak with respect to synchrotron and IC.  A key prediction of our scenario with $B \\propto \\Sigma_g^{0.7-0.8}$ is that the variations in $L_{\\rm TIR}^{\\prime}/L_{\\rm radio}^{\\prime}$ will be correlated with scale height at fixed $\\Sigma_{\\rm SFR}$, since the radio-excess in our models is a direct consequence of the large CR scale height and the small bremsstrahlung and ionization losses it causes.  Radio-excess (low $q$) puffy starbursts will have steeper radio spectral slopes (bigger $\\alpha$), larger velocity dispersions $\\sigma$ compared to their rotation speeds $v_{\\rm circ}$, and possibly moderately cooler dust temperatures (smaller $T_{\\rm dust}$).  \\item As previously expected, radio emission can be a poor tracer of star formation in low surface density galaxies, because of electron escape and IC losses off the CMB.  For our preferred $B \\propto \\Sigma_g^{0.7-0.8}$ scaling, radio emission overestimates star-formation rate by a factor of $2 - 4$ in puffy starbursts.  Star-formation rate is underestimated by synchrotron radio emission with the $B \\propto \\rho^{0.5 - 0.6}$ scaling.  \\item While SMGs may be individually radio bright compared to the local FRC, they contribute a relatively small fraction of the Cosmic Infrared Background and the total star-formation luminosity of the Universe.  This means that they enhance the star-formation radio background by $\\la 50\\%$, and possibly around $\\sim 10\\%$, with respect to a naive application of the $z = 0$ FRC. \\end{enumerate}  As in LTQ, we did not exactly match the observed radio spectral slopes of galaxies, with $\\alpha \\approx 0.9 - 1.0$ in normal galaxies and $\\alpha \\approx 0.4 - 0.6$ in compact starbursts.  This will have a slight effect on the k-correction.  An error of $0.25$ in $\\alpha$ should only affect inferred radio luminosity by 30\\% at $z = 2$ and 60\\% at $z = 5$.  Nevertheless, the prediction of steeper radio spectra in puffy starbursts with respect to compact starbursts at all relevant $z$ should be robust.  Our explanation for the small $L_{\\rm TIR}^{\\prime}/L_{\\rm radio}^{\\prime}$ ratio in submillimeter galaxies as a breakdown of the high-$\\Sigma_g$ conspiracy is based purely on the steady-state spatially-averaged synchrotron emission, but the details of the FIR emission may also matter.  Throughout this paper, we have simply assumed that the bolometric FIR luminosity could be correctly inferred from observations, and have assumed the same UV opacity for all galaxies at all redshifts.  The total FIR emission is also likely to depend on the metallicity, and may be lower at high $z$ for the lowest surface density galaxies.  The exact far-infrared SED is important in determining the FIR emission when observations have only been made at only a few wavelengths.  The presence of AGNs, a different IMF at high $z$, and selection biases may also affect the inferred $q$ of SMGs.  We did not include the effects of galaxy evolution on the CR spectrum in our models.  We argued in LTQ that it should not matter for quiescent spirals or for extreme starbursts, because the CR lifetime is much shorter than the time dependence of stellar populations.  However, galaxy evolution may play a role in weaker starbursts \\citep{Lisenfeld96b} and in post-starburst galaxies \\citep{Bressan02}.  Studies of merging normal galaxies and galaxies in clusters have indeed found that they are radio bright with respect to the FRC, possibly because of compression of magnetic fields or shock acceleration \\citep{Gavazzi91,Miller01,Murphy09b}.  We also assumed that the magnetic field strength at a given density does not depend on redshift.  It is not entirely clear how long normal galaxies take to build up their magnetic fields, or even what process is at work \\citep[see the reviews in ][]{Widrow02,Kulsrud08}, though there are theoretical mechanisms that can rapidly generate strong magnetic fields.  Studies at $z \\approx 2$ indicate that normal Milky Way-like galaxies had magnetic fields with similar strengths to the present \\citep{Kronberg08,Bernet08,Wolfe08}.  At the very highest redshifts, magnetic field strengths might be weaker, because the seed fields were essentially zero compared to the present strengths.  Starbursts also may build their magnetic fields up in much shorter times than normal galaxies, and through a different process than normal galaxies \\citep{Thompson09}.  Finally, we have used one-zone models, which are appropriate if the CRs sample all of the gas phases in each galaxy's ISM.  However, the ISM is known to be clumpy in the Milky Way, in compact starbursts like Arp 220 \\citep[e.g.,][]{Greve09}, and even in the SMGs themselves \\citep{Tacconi06}.  A full understanding of the FRC will probably require models that take into account the inhomogeneity of star-forming galaxies."
    },
    "0910/0910.3960_arXiv.txt": {
        "abstract": "Galactic bars are the most important driver of secular evolution in galaxies. They can efficiently drive gas into the central kiloparsec of galaxies, thus feed circumnuclear starbursts, and possibly help to fuel AGN.  The connection between bars and AGN activities has been actively debated in the past two decades. Previous work used fairly small samples and often lacked a proper control sample. They reported conflicting results on the correlation between bars and AGN activity. Here we revisit the bar-AGN and bar-starburst connections using the analysis of bars in a large sample of about 2000 SDSS disk galaxies \\citep*{bar_etal_08}. We find that AGN and star-forming galaxies have similar optical bar fractions, 47\\% and 50\\%, respectively. Both bar fractions are higher than that in inactive galaxies  (29\\%). We discuss the implications of the study on the relationship between host galaxies and their central activities. ",
        "introduction": "Large-scale bars are very common in disk galaxies. At near infrared (NIR) wavelengths, the optical bar fraction averaged across different Hubble types  is $\\sim$~60\\%  from quantitative bar identification methods, such as the ellipse fitting and structural decomposition \\citep*[e.g.][]{lau_etal_04a, mar_jog_07, men_etal_07, wei_etal_09} , and is $\\sim$~72\\% from visual classification \\citep{esk_etal_00}. At optical wavelengths,  quantitative methods yield an average optical bar fraction of  45\\% to 52\\%  \\citep*{mar_jog_07, bar_etal_08, agu_etal_09}, while visual classification yields  $\\sim$~60\\% \\citep{devauc_63}. The optical bar fraction is somewhat lower than the NIR fraction due to the obscuration of bars  by dust lanes and star formation along the bar.  Several studies have now moved beyond considering only the bar fraction averaged over many Hubble types. They investigated how the optical bar fraction varies with the Hubble types or host properties. The optical bar fraction rises in galaxies with low Bulge-to-Disk ratio or/and high luminosity \\citep{odewah_96, bar_etal_08, bar_etal_09, agu_etal_09, mar_etal_09}.  The non-axisymmetric stellar bar can drive gas from the outer disk to the central kiloparsec, where they trigger star formation. This is supported by several observations, which show that barred galaxies host high gas densities and circumnuclear starbursts \\citep*[e.g.][]{jog_etal_05} and that the central gas concentration is larger in barred than unbarred galaxies \\citep*[e.g.][]{sak_etal_99, she_etal_04}. Due to their efficiency in driving gas inflows in the disks, bars are strong candidates for the triggering of nuclear activities.  There is strong evidence for a connection between large-scale bars and circumnuclear starbursts. Barred galaxies show enhanced radio continuum and infrared emissions compared with unbarred ones \\citep[e.g.][]{hummel_81, haw_etal_86}, and starburst galaxies tend to be more barred compared with the non-starburst galaxies \\citep*[e.g.][]{arsena_89, hua_etal_96, ho_etal_97, hun_mal_99}. Bars can also set up resonances, such as the inner Lindblad resonances (ILRs), which can prevent the gas from going further in. Therefore, gas often builds up on a ring (a few hundred parsecs in radius) and the circumnuclear starbursts can occur there \\citep[e.g.][]{jog_etal_05}.  The connection between bars and AGNs is less clear (see \\citealt{jogee_06} for a review). Over the last two decades a large number of studies were carried out to identify if such a correlation exists. Most of them compared the bar fraction of the AGN sample with that of a control sample of inactive galaxies. The results are controversial. For example, studies like \\citet{ho_etal_97}, \\citet{hun_mal_99}, \\citet{mul_reg_97}, and \\citet{mar_etal_03} found no excess of bars in Seyfert galaxies, while \\citet*{kna_etal_00}, \\citet{hao_lai_etal_02}, and \\citet*{lau_etal_04b} reported a higher fraction of bars in Seyferts.   Previous studies are limited in several aspects. Firstly, the sizes of the samples are often small, including only a few tens of AGNs and control galaxies. In some cases, the control sample was not well matched to the active sample. Secondly, identifications method for bars are not always consistent across samples. Thirdly, the spectral classifications of galaxies as AGNs or inactive galaxies are significantly inconsistent. Most studies adopted the galaxy classifications in NASA/IPAC NED, which could be done by different people using different criteria. Our study tries to overcome some of  these limitations by systematically investigating the optical bar fraction of AGNs and non-AGNs in a large number of galaxies from the Sloan Digital Sky Survey (SDSS), using matched active and control samples, and the same consistent quantitative method for identifying bars across samples. ",
        "conclusions": "With the classification and structural information of $\\sim 2000$ disk galaxies from the SDSS \\citep{bar_etal_08}, we study the optical bar fraction of AGNs, star-forming, and inactive galaxies from the sample. We find that the optical bar fraction of the AGNs is 47\\%, similar to the optical bar fraction of the star-forming galaxies (50\\%). Both are higher than the optical bar fraction of the inactive galaxies (29\\%). This suggests that accurate and consisitent spectral classification is important in evaluating whether there is an excess of bars in AGNs, and could be the reason for controversial results reported in previous studies on the issue. Our study has several imporvements compared to previous ones. The size of our sample is large. We have consistent SDSS spectra for all our galaxies, therefore, we can obtain accurate and consistent spectral classifications for the sample. In addition, the SDSS spectra are taken with the 3\\arcsec aperture, which corresponds to 606 pc to 1.78 kpc in the redshift range of $0.01<z<0.03$. This scale matches well with the typical circumnuclear region of a galaxy, therefore, the spectra are perfect at probing the circumnuclear stubursts.  We find an excess of the optical bar fraction of star-forming galaxies and AGNs compared to the inactive galaxies. This suggests that large-scale primary bars drive gas to inner kpc where they pile up near the ILRs, fueling  circumnuclear starbursts. The gas pile up is in someway also related with the increase of the AGN activity. But to feed the AGN directly, the gas in the inner kpc still has to reduce its angular momentum by several orders of magnitude  (see Figure 3 in \\citealt{jogee_06}) and a secondary mechanism is then needed to drive the gas further in  (e.g., nuclear bars, dynamical friction). The latter may not be necessarily coupled to the primary bar. This agrees with our result that we do not observe an excess of bar fraction of AGNs compared to the star-forming galaxies. Furthermore, we find no evidence of bar weakening by AGNs. This agrees with previous theoretical expectations \\citep{she_sel_04, ath_etal_05}.  There is one caveat of our study. Our bar identification is based on the optical data instead of NIR, therefor our bar fraction is in general lower than the typical NIR bar fraction. However, we do not expect our comparative results to change, unless the obscuration of bars by dust and star formation have preferential effects."
    },
    "0910/0910.3367_arXiv.txt": {
        "abstract": "Borexino is a solar neutrino experiment running at the Laboratori Nazionali del Gran Sasso, Italy. The radioactive background levels in the liquid scintillator target meet or even exceed design goals, opening unanticipated opportunities. The main results, so far, are the measurement of the $^7$Be solar neutrino flux (the first ever done) and the measurement of the $^8$B neutrino flux performed with electron energy threshold of 2.8 MeV. The short and medium term perspectives are summarized in the conclusions. ",
        "introduction": "Borexino \\cite{BxNim} detects solar neutrinos via their elastic scattering on the electrons of an ultra-pure liquid scintillator target. The main physics goal  is the measurement of the flux and of the energy spectrum of solar neutrinos with sub-MeV or few MeV energy. Other goals include geoneutrinos detection, super-nova neutrinos detection and the search for very rare decays \\cite{CTF1,CTF2,CTF3,CTF4}.  Flavor oscillations of solar neutrinos with MSW effect \\cite{MSW} have been well established by radiochemical experiments \\cite{RadioChemical} and by water Cerenkov detectors \\cite{Cerenkov}. The range of the parameters describing the oscillation phenomenon  has been constrained by SNO and KamLand\\cite{Kamland} to lie in the so called LMA (Large Mixing Angle) region of the plane $ \\theta_{12}$ $\\Delta m^2_{12}$ ($tan^2 (2\\theta_{12}) =0.47^{ +0.06}_{-0.05} $ and $\\Delta m^2_{12} 7.59^{+0.21}_{-0.21} \\cdot 10^{-5} eV^2$).  The main prediction of the LMA-MSW are the energy dependence  of the neutrino survival probability $P_{ee}$ and the lack of day-night asymmetry. The $P_{ee}$ decreases with increasing energy. Matter effects dominate at energies above 3 MeV while are absent below 1 MeV. The region in between is called the transition region. While the LMA-MSW predicts a well defined shape for the  $P_{ee}$ function in the transition region, current experimental data do not constraint it at all, and some theoretical models, including non standard interactions, predict survival probability curves  with different shape \\cite{NotStandard}.  Borexino is presently the only running experiment that measured the signal rate due to the 0.862 MeV $^7$Be neutrinos \\cite{BxBe}\\cite{BxBe2} and the one due to $^8B$ \\cite{BxB8} with  a energy threshold (2.8 MeV), lower than any previous experiment. Besides, a measurement of a null day night difference of the $^7$Be flux provides a further new confirmation of the LMA-MSW solution. Interest of this measurement is also related to the possibility to accommodate  quite large effects in some alternative oscillation scenario based on the mass varying model \\cite{Holanda}.  The neutrino interaction rate measured by Borexino depends on  the solar neutrino flux and on the oscillation parameters. A scientific debate about high and low metallicity solar models and the related flux calculations \\cite{carlos1} is in progress. The relevance of the measurements of the various components of the solar neutrino measurements in Borexino is then twofold: from one side they can increase the confidence in the oscillation scenario and from the other side, assuming the knowledge of the oscillation parameters, they provide a measurement of the absolute solar neutrino flux. The precision of the actual data about solar neutrinos  does not allow to distinguish between high and low metallicity models \\cite{carlos1} but future high precision measurements of the $^7$Be might give useful constraints.  Besides, the experimental determination of the flux of the CNO components is of strategical importance being the CNO flux prediction different by more than $30 \\%$ in the two classes of models. ",
        "conclusions": "Borexino has completed the first real time measurements of $^7$Be and the low threshold $^8$B signal rate. A preliminary result about the day night asymmetry has also been obtained. All results confirm the current LMA-MSW scenario. Annual modulation analysis is underway.  A calibration campaign is in progress. The results are expected to contribute to a significant reduction of the errors on the fiducial volume and on the detector response function, yielding a $^7$Be signal rate measurement with few percent accuracy and of the reduction of the systematic errors of the $^8B$ measurement.  The possibility to purify the scintillator to reduce the background due to $^{85}Kr$ and to $^{210}Bi$ is under consideration. If successful, Borexino might attempt the direct detection of pep and pp, and CNO neutrinos."
    },
    "0910/0910.3684_arXiv.txt": {
        "abstract": "I present the results of Monte-Carlo orbital simulations of Galactic Neutron Stars (NSs). The simulations take into account the up-to-date observed NS space and velocity distributions at birth, and account for their formation rate. I simulate two populations of NSs. Objects in the first population were born in the Galactic disk at a constant rate, in the past 12\\,Gyr. Those in the second population were formed simultaneously 12\\,Gyr ago in the Galactic bulge. I assume that the NSs born in the Galactic disk comprise 40\\% of the total NS population. Since the initial velocity distribution of NSs is not well known, I run two sets of simulations, each containing $3\\times10^{6}$ simulated NSs. One set utilizes a bimodal initial velocity distribution and the other a unimodal initial velocity distribution, both are advocated based on pulsars observations. In light of recent observational results, I discuss the effect of dynamical heating by Galactic structure on NS space and velocity distributions and show it can be neglected. I present catalogue of simulated NS space and velocity vectors in the current epoch, and catalogue of  positions, distances and proper motions of simulated NSs, relative to the Sun. Assuming there are $10^{9}$ NSs in the Galaxy, I find that in the solar neighborhood the density of NSs is about $2-4\\times10^{-4}$\\,pc$^{-3}$, and their scale height is about $0.3-0.6$\\,kpc (depending on the adopted initial velocity distribution). These catalogue can be used to test the hypothesis that some radio transients are related to these objects. ",
        "introduction": "\\label{sec:Introduction}  The present day space and velocity distributions of Galactic isolated old Neutron Stars (NS), are the subject of many studies. Among the reasons for deriving the positional and kinematical properties of isolated old NSs was the suggestion that Gamma-Ray Bursts (GRBs) may originate from Galactic NSs (e.g., Mazets et al. 1980). This was refuted\\footnote{At least in the sense that the majority of GRBs are extragalactic; see Kasliwal et al. (2008).} by the homogeneous sky distribution of GRBs (e.g., Meegan et al. 1992), and later on by the discovery of their cosmic origin (e.g., Metzger et al. 1997; van Paradijs et al. 1997).  Another exciting possibility, that re-ignited these efforts, was to detect isolated old NSs in soft X-ray radiation that may be emitted as they slowly accrete matter from the interstellar medium (ISM; Ostriker, Rees \\& Silk 1970; Shvartsman 1971). Treves \\& Colpi (1991), and Blaes \\& Madau (1993) predicted that $\\sim10^{3}$--$10^{4}$ isolated old NS accreting from the ISM would be detected by the {\\it ROSAT} Position Sensitive Proportional Counters (PSPC) all sky survey (Voges et al. 1999). Although intensive searches for these objects were carried out (e.g., Motch et al. 1997; Maoz, Ofek, \\& Shemi 1997; Haberl, Motch, \\& Pietsch 1998; Rutledge et al. 2003;  Ag{\\\"u}eros et al. 2006), only a handful of candidates were found. However, these are presumably young cooling NSs, rather than old NSs whose luminosity is dominated by accretion from the ISM (e.g., Neuh\\\"{a}user \\& Tr\\\"{u}mper 1999; Popov et al. 2000; Treves et al. 2001). Apparently, the reasons for the rareness of these objects in the {\\it ROSAT} source catalog is that their typical velocities were underestimated by earlier studies (i.e., Narayan \\& Ostriker 1990). It is also possible that their accretion rate is below the Bondi-Hoyle rate\\footnote{Bondi \\& Hoyle (1944).} (e.g., Colpi et al. 1998; Perna et al. 2003; Toropina et al. 2001, 2003, 2005; see however Arons \\& Lea 1976,1980), or that these NSs are in an ejection stage (e.g., Colpi et al. 1998; Livio et al. 1998; Popov \\& Prokhorov 2000).  During the last several years, new classes of radio transients were found (e.g., Bower et al. 2007; Matsumura et al. 2007; Niinuma et al. 2007; Kida et al. 2008; see also Levinson et al. 2002; Gal-Yam et al. 2006). Bower et al. (2007) found several examples of these transients in a single small field of view, and showed that they have time scales above 20\\,min but below seven days, and they lack any quiescent X-ray, visible-light, near-Infrared (IR), and radio counterparts. Recently, Ofek et al. (2009) suggested that these transient events may be associated with Galactic isolated old NSs. However, testing this hypothesis requires knowledge of the NSs space distribution.   There are two approaches in the literature for calculating the theoretical space and velocity distributions of old NSs. Paczynski (1990), Hartmann, Woosley, \\& Epstein (1990), Blaes \\& Rajagopal (1991), Blaes \\& Madau (1993), and Posselt et al. (2008) carried out Monte-Carlo simulations of NS orbits. In such simulations, the positions and velocities of NS at birth are integrated assuming some non-evolving Galactic potential. A second approach is to use some sort of semi-analytic approximation in order to estimate the ``final'' NS space and velocity distributions. Frei, Huang, \\& Paczynski (1992) calculated the final vertical density and velocity distributions of NSs, assuming that the gravitational potential is a function only of the height, $z$, above the Galactic plane. Using the epicyclic approximation, Blaes \\& Rajagopal (1991) developed a technique that allows calculation of the full three dimensional velocity distribution of NSs. However, they showed that this method is not adequate for fast moving objects, which constitute the majority of the NS population. Blaes \\& Madau (1993) used the thin-disk approximation to calculate the space and velocity distributions of NSs. In this prescription, the radial motion of NSs are controlled by the Galactic potential in the Galactic disk, regardless of the vertical height, $z$, of the NS above/below the Galactic plane. Finally, another solution was presented by Prokhorov \\& Postnov (1994), who assume that the ergodic hypothesis is correct (see discussion in Binney \\& Tremaine 1987).   Madau \\& Blaes (1994) noted that all these approaches neglect dynamical heating of NSs due to encounters with giant molecular clouds, spiral arms, and stellar ``collisions'' (e.g., Kamahori \\& Fujimoto 1986, 1987; Barbanis \\& Woltjer 1967; Carlberg \\& Sellwood 1985; Jenkins \\& Binney 1990). They crudely estimated the order of magnitude of this effect by applying the force-free diffusion equation to the vertical height, $z$, and NS speed distributions. They found that, at the solar neighborhood, dynamical heating may decrease the local density of NSs by a factor of $\\sim1.5$ relative to a non-heated population.    In this paper I present the results of Monte-Carlo orbital simulations of Galactic NSs. These simulations improve upon past efforts (e.g., Paczynski 1990; Hartmann, Epstein, \\& Woosley 1990, Blaes \\& Rajagopal 1991; Blaes \\& Madau 1993; Popov et al. 2005; and Posselt et al. 2008) in several aspects. First, I use up-to-date space and velocity distributions of NSs at birth from Arzoumanian et al. (2002) and Faucher-Gigu\\`{e}re \\& Kaspi (2006). Second, I assume a birth rate of NSs along the Galaxy life time, instead of assuming that all the NSs were born about 10~Gyr ago. Third, I generate a large sample of simulated NSs, which is about one to two orders of magnitude larger than those presented by previous efforts.   The structure of this paper is as follows: In \\S\\ref{sec:ModelIng} I present the model ingredients for the Monte-Carlo simulations. In \\S\\ref{sec:DynHeat} I discuss dynamical heating and show that it can be neglected. The Monte-Carlo simulations are described in \\S\\ref{sec:MC}, and the catalogue of simulated NSs are presented in \\S\\ref{sec:Cat}. The results of these simulations are presented in \\S\\ref{sec:NSres}, and in \\S\\ref{sec:Disc} I summarize the results and compare them with some of the previous efforts. ",
        "conclusions": "\\label{sec:Disc}  The Milky Way's NSs content, and in particular the NS space and velocity distributions are important for searching nearby isolated old NSs, and studying any ongoing activity from such objects. As a tool for such studies, I present a mock catalogue of simulated isolated old NSs spatial positions and velocities, at the current epoch.  The catalogue were constructed by integrating the equations of motion of simulated NSs in the Galactic potential. The simulations include two populations of NSs, one in which the NSs were born in the Galactic bulge about 12\\,Gyr ago, and the second population in which NSs are being born at the Galactic disk at a constant rate, starting 12\\,Gyr ago. The combined NS population assumes that $60\\%$ of the NSs originated in the bulge and the rest in the disk. Although we do~not know what is the exact number of Galactic NSs, and what was their position-dependent birth rate, these two populations provide a wide range of initial conditions.  I generated two catalogue of simulated NSs. Each catalog contains $3\\times10^{6}$ objects. The two catalogue utilize different initial velocity distributions of NSs. One catalog (Table~\\ref{tab:MC_v2g}) uses the initial velocity distribution of Arzoumanian et al. (2002), while the other (Table~\\ref{tab:MC_v1e}) uses the initial velocity distribution of Faucher-Gigu\\`{e}re \\& Kaspi (2006). Also derived are catalogue of simulated NS positions and proper motions with respect to an observer found at the solar circle (tables~\\ref{tab:SunCentric2} and \\ref{tab:SunCentric1}).  The space distribution at the current epoch obtained by the two initial velocity distributions implemented here, are somewhat different. For example, I find that the resulting sky surface density, based on the two different initial velocity distributions, differs by up to $65\\%$. The main differences between the two velocity distributions are in the low- and high-tails of the NSs velocity distribution (see Fig.~\\ref{fig:InitVelDist_Compare}).    \\begin{deluxetable}{lccl} \\tablecolumns{4} \\tablewidth{0pt} \\tablecaption{Comparison with previous works} \\tablehead{ \\colhead{$\\rho/N_{9}$} & \\colhead{$Z_{1/2}$} & \\colhead{mode$(p[v])$\\tablenotemark{a}} & \\colhead{Reference} \\\\ \\colhead{pc$^{-3}$} & \\colhead{kpc} & \\colhead{km\\,s$^{-1}$} & \\colhead{} } \\startdata $2.4\\times10^{-4}$  &0.42  & 240      & This Work (disk$+$bulge) A2002\\tablenotemark{b} \\\\ $4.0\\times10^{-4}$  &0.20  & 200      & This Work (disk$+$bulge) FK2006\\tablenotemark{c} \\\\ $1.4\\times10^{-3}$  &0.20  &          & Paczynski (1990) \\\\ $7.5\\times10^{-4}$  &0.27  &  43      & Blaes \\& Madau (1993) \\\\ $5.3\\times10^{-4}$  &0.50  &  69      & Madau \\& Blaes (1994) \\\\ $4\\times10^{-4}$    &      & 140      & Perna et al. (2003) \\\\ $5\\times10^{-4}$    &0.4   &          & Popov et al. (2005) \\enddata \\tablenotetext{a}{Mode is the most probable value of the distribution, and the speeds are measured relative to an inertial reference frame.} \\tablenotetext{b}{Using the bimodal initial velocity distribution of Arzoumanian et al. (2002).} \\tablenotetext{c}{Using the unimodal double-sided exponential initial velocity distribution of Faucher-Gigu\\`{e}re \\& Kaspi (2006).} \\tablecomments{The density, $\\rho$, at the solar neighborhood is calculated assuming there are $10^{9}$ NSs in the Galaxy. $Z_{1/2}$ is the height above the Galactic plane, in which the NS density drops to $1/2$ its value on the Galactic plane, calculated at the solar circle. The density from Madau \\& Blaes (1994) is calculated assuming that diffusion operats for 5\\,Gyr (see \\S\\ref{sec:DynHeat}). The initial velocity distribution used by Paczynski (1990), Blaes \\& Madau (1993), and Madau \\& Blaes (1994) is from Narayan \\& Ostriker (1990). The initial velocity distribution in Perna et al. (2003) is from Cordes \\& Chernoff (1998). } \\label{tab:Comp} \\end{deluxetable} The space and velocity distributions of Galactic NSs were estimated by several past works. In Table~\\ref{tab:Comp}, I compare some of the basic results of the simulations presented here (e.g., local density of NSs; scale height), with the ones obtained by previous efforts."
    },
    "0910/0910.3017_arXiv.txt": {
        "abstract": "Probabilities of collisions of migrating small bodies and dust particles produced by these bodies with planets were studied. Various Jupiter-family comets, Halley-type comets, long-period comets, trans-Neptunian objects, and asteroids were considered. The total probability of collisions of any considered body or particle with all planets did not exceed 0.2. The amount of water delivered from outside of Jupiter's orbit to the Earth during the formation of the giant planets could exceed the amount of water in Earth's oceans. The ratio of the mass of water delivered to a planet by Jupiter-family comets or Halley-type comets to the mass of the planet can be greater for Mars, Venus, and Mercury, than that for Earth. ",
        "introduction": "In the present paper we discuss the probabilities of collisions of migrating small bodies and dust particles produced by these bodies with planets. Ipatov \\& Mather (2003, 2004a-b, 2006, 2007), Ipatov et al. (2004),  and several other authors cited in the above papers studied the probabilities with Earth, Venus, and Mars. In our above papers (which can be found on astro-ph and http://faculty.cua.edu/ipatov/list-publications.htm), the probabilities of the collisions were calculated based on the time variations of orbits of bodies and particles during their dynamical lifetimes (until all bodies or particles reached 2000 AU from the Sun or collided with the Sun). Using  the same time variations, below we consider also the probabilities of collisions with the giant planets and Mercury.  The orbital evolution of $>$30,000 bodies with initial orbits close to those of Jupiter-family comets (JFCs), Halley-type comets, long-period comets, trans-Neptunian objects, and asteroids in the resonances 3/1 and 5/2 with Jupiter, and also of $>$20,000 dust particles produced by these small bodies was integrated  during their dynamical lifetimes. We considered the gravitational influence of planets, but omitted the influence of Mercury (exclusive for Comet 2P/Encke). In about a half of calculations of migration of bodies, we used the method by Bulirsh-Stoer (1966) (BULSTO code), and in other runs we used a symplectic method (RMVS3 code). The integration package of Levison \\& Duncan (1994) was used. For dust particles, only the BULSTO code was used, and the gravitational influence of all planets, the Poynting-Robertson drag, radiation pressure, and solar wind drag were taken into account. The ratio $\\beta$ between the radiation pressure force and the gravitational force varied from $\\le$0.0004 to 0.4. For silicates, such values of $\\beta$ correspond to particle diameters $d$ between $\\ge$1000 and 1 microns; $d$ is proportional to 1/$\\beta$.   In our calculations, planets were considered as material points, so literal collisions did not occur. However, using the algorithm suggested by Ipatov (1988) with the correction that takes into account a different velocity at different parts of the orbit (Ipatov \\& Mather 2003), and based on all orbital elements sampled with a 10-500 yr step, we calculated the mean probability $P$ of collisions of migrating objects with a planet. The step could be different for different typical dynamical lifetimes of particles. We define $P$ as $P_{\\Sigma}/N$, where $P_{\\Sigma}$ is the probability of a collision of all $N$ objects  with a planet. The probabilities of collisions of bodies and particles at different $\\beta$ with planets % are presented in Fig. 1. These probabilities do not take into account the destruction of particles in collisions and sublimation, which can be more important for small particles. Our runs were made mainly for direct modelling of collisions with the Sun, but Ipatov \\& Mather (2007) obtained that the mean probabilities of collisions of considered bodies with planets, lifetimes of the bodies that spent millions of years in Earth-crossing orbits, and other obtained results were practically the same if we consider that bodies disappear when  perihelion distance becomes less than the radius of the Sun or even several such radii. ",
        "conclusions": ""
    },
    "0910/0910.0318.txt": {
        "abstract": "Collimated supersonic flows in laboratory experiments behave in a similar manner to astrophysical jets provided that radiation, viscosity, and thermal conductivity are unimportant in the laboratory jets, and that the experimental and astrophysical jets share similar dimensionless parameters such as the Mach number and the ratio of the density between the jet and the ambient medium.  When these conditions apply, laboratory jets provide a means to study their astrophysical counterparts for a variety of initial conditions, arbitrary viewing angles, and different times, attributes especially helpful for interpreting astronomical images where the viewing angle and initial conditions are fixed and the time domain is limited. Experiments are also a powerful way to test numerical fluid codes in a parameter range where the codes must perform well. In this paper we combine images from a series of laboratory experiments of deflected supersonic jets with numerical simulations and new spectral observations of an astrophysical example, the young stellar jet HH~110. The experiments provide key insights into how deflected jets evolve in 3-D, particularly within working surfaces where multiple subsonic shells and filaments form, and along the interface where shocked jet material penetrates into and destroys the obstacle along its path. The experiments also underscore the importance of the viewing angle in determining what an observer will see. The simulations match the experiments so well that we can use the simulated velocity maps to compare the dynamics in the experiment with those implied by the astronomical spectra. The experiments support a model where the observed shock structures in HH 110 form as a result of a pulsed driving source rather than from weak shocks that may arise in the supersonic shear layer between the Mach disk and bow shock of the jet's working surface. ",
        "introduction": "Collimated supersonic jets originate from a variety of astronomical sources, including active galactic nuclei \\citep{agnref}, several kinds of interacting binaries \\citep{binaryjetref}, young stars \\citep{ysojetref}, and even planetary nebulae \\citep{pnjetref}. Most current jet research focuses on how accretion disks accelerate and collimate jets, or on understanding the dynamics of the jet as it generates shocks along its beam and in the surrounding medium. Both areas of research have broad implications for astrophysics. Models of accretion disks typically employ magnetized jets to remove the angular momentum of the accreting material. The distribution and transport of the angular momentum in an accretion disk affects its mass accretion rate, mixing, temperature profile, and density structure, and in the case of young stars, also helps to define the characteristics of the protoplanetary disk that remains after accretion ceases.  At larger distances from the source, shock waves in jets clear material from the surrounding medium, provide insights into the nature of density and velocity perturbations in the flow, and enable dynamical studies of mixing, turbulence and shear.  Jets from young stars are particularly good testbeds for investigating all aspects of the physics within collimated supersonic flows \\citep[see][for a review]{ray07}. Shock velocities within stellar jets are low enough that the gas cools by radiating emission lines rather than by expanding. Relative fluxes of the emission lines determine the density, temperature, and ionization of the postshock gas, while the observed Doppler shifts and emission line profiles define the radial velocities and nonthermal motions within the jet \\citep[see][for a review]{hartigan08}. Moreover, many stellar jets are located relatively close to the Earth, so that one can observe proper motions in the plane of the sky from observations separated by several years \\citep{heathcote92}. Combining this information with radial velocity measurements gives the orientation of the flow to the line of sight. Using the Hubble Space Telescope, one can observe morphological changes of knots within jets, and follow how these changes evolve in real time \\citep{hartigan01}.  Results from these studies show that internal shock waves, driven by velocity variations in the flow, sweep material in jets into a series of dense knots. Typical internal shock velocities are $\\sim$ 40 km$\\,$s$^{-1}$, or $\\sim$ 20\\%\\ of the flow speed. In several cases it is easy to identify both the bow shock and the Mach disk from emission line images. Because stellar jets are mostly neutral, strong H$\\alpha$ emission occurs at the shock front where neutral H is collisionally excited \\citep{heathcote96}. Forbidden line radiation occurs in a spatially-extended cooling zone in the postshock material. Temperatures immediately behind the internal shocks can exceed $10^5$K, but the gas in the forbidden-line-emitting cooling zone is typically 8000~K.  Young stars show a strong correlation between accretion and outflow, leading to the idea that accretion disks power the outflows \\citep{heg95,cabrit07}. Most current models use magnetic fields in the disk to launch a fraction of the accreting material from the disk into a collimated magnetized jet \\citep{ferreira06}. Stellar jets often precess, and there is some evidence that they rotate \\citep{ray07}, although rotation signatures are difficult to measure because the rotational velocities are typically only a few percent of the flow speeds and precession can mimic rotational signatures \\citep{cerquieriaref}.  In all cases proper motion measurements show that jets move radially away from the source.  Jets show no dynamical evidence for kink instabilities, and in fact while magnetic fields may dominate in the acceleration regions of jets, they appear to play a minor role in the dynamics at the distances where most jet knots are observed \\citep{hartigan07}. In the cooling zones behind the shock waves the plasma $\\beta$ can drop below unity, so that magnetic fields dominate thermal pressure in those areas. However, the magnetic pressure is small compared with the ram pressure of the jet.  While more unusual than internal working surfaces produced by velocity perturbations, shocks also occur when jets collide with dense obstacles such as a molecular cloud. When the obstacle is smaller than the jet radius, it becomes entrained by the jet, and a reverse bow shock or `cloudlet shock' forms around the obstacle \\citep{schwartz78,l1551}. Alternatively, when a large obstacle like a molecular cloud deflects the jet, a quasi-stationary deflection shock at the impact point forms, followed by a spray of shocked jet material downstream. The classic example of such a jet is HH 110 \\citep{riera03,lopez05}.  Though the observations summarized above provide a great deal of information about stellar jets, several important questions remain unanswered. The basic mechanism by which disks load material onto field lines (assuming the MHD disk scenario is correct), and the overall geometry of this wind is unknown, and the roles of reconnection and ambipolar diffusion in heating the jet close to the star are unclear. At larger distances, the magnetic geometry and its importance in shaping the internal working surfaces is poorly-constrained, as are the time scales and spatial scales associated with mixing in supersonic shear layers and working surfaces. The inherently clumpy nature of jets also affects the flow dynamics and observed properties of jets in uncertain ways, and the degree to which fragmentation and turbulence influence the morphologies of jets is unknown.  Developing laboratory analogs of stellar jets could help significantly in addressing the questions above. Observations of a specific astronomical jet are restricted to a small range of times and to a particular observing angle, while laboratory experiments have no such restrictions. In principle, one could explore a wide range of initial collimations, velocity and density structures within the jet, as well as densities, geometries, and magnetic field configurations with laboratory experiments.  The experiments also provide a powerful and flexible way to test 3-D numerical fluid codes, and to investigate how real flows develop complex morphologies in 3-D.  The challenge is to design an experiment that is relevant to the astrophysical case of interest. Laboratory experiments differ by 15 $-$ 20 orders of magnitude in size, density and timescale from stellar jets, but because the Euler equations that govern fluid dynamics involve only three variables, time, density, and velocity, the behavior of the fluid is determined primarily by dimensionless numbers such as the Mach number (supersonic or subsonic), Reynolds number (viscous or inertial), and Peclet number (importance of thermal conduction). If the experiments behave as a fluid and have similar dimensionless fluid numbers as those of stellar jets, then the experiments should scale well to the astrophysical case \\citep{ryutov99}. Other parameters, such as magnetic fields and radiational cooling are more difficult to match, and it is impossible to study the non-LTE excitation physics in the lab because the critical density for collisional deexcitation is not scalable. In any case, the materials are markedly different between the experiments and stellar jets, so it is not possible to study emission line ratios in any meaningful way with current laboratory capabilities. Hence, at present the main utility of laboratory experiments of jets is to clarify how complex supersonic flows evolve with time.  Laboratory work relevant to stellar jets is an emerging area of research, and several papers have appeared recently which address various aspects of supersonic flows in the appropriate regime \\cite[see][for a review]{remington06}. \\citet{hansen07} observed how a strong planar shock wave disrupts a spherical obstacle and tested numerical models of the process, and \\citet{loupias} generated a laboratory jet with a Mach number similar to that of a stellar jet. In a different approach, \\citet{lebedev04} and \\citet{ampleford07} used a conical array of wires at the Magpie facility to drive a magnetized jet, and explored the geometry of the deflection shock formed as a jet impacts a crosswind. Laboratory experiments have also recently studied the physics associated with instabilities along supernova blast waves \\citep{snrref,drake09}, and the dynamics within supernovae explosions \\citep{snrcoreref}.  In this paper we present the results of a suite of experiments which deflect a supersonic jet from a spherical obstacle, where the dimensionless fluid parameters are similar to those present in stellar jets. In section 2, we describe our experimental design, consider how the experiments scale to astrophysical jets, demonstrate that the experiments are reproducible, and report how the observed flows change as one varies several parameters, including the distance between the axis of the jet and the obstacle, the time delay, the density probed with different backlighters, and the viewing angle. The numerical work is summarized in section 3. Detailed calculations with the 3-D RAGE code reproduce all of the major observed morphologies well. In section 4, we present new high-resolution optical spectra of the shocked wake of the HH 110 protostellar jet, and a new wide-field H$_2$ image of the region. These observations quantify how the internal dynamics of the gas behave as material flows away from the deflection shock and show how the jet entrains material from the molecular cloud core. Finally, in section 5 we consider how the experiments and the simulations from RAGE provide new insights into the internal dynamics of deflected supersonic jets, especially in the regions of the working surface and at the interface where jet material entrains and accelerates the obstacle. ",
        "conclusions": "%6  The combination of experimental, numerical and astronomical observational data from this study demonstrates the potential of the emerging field of laboratory astrophysics. In this paper we have studied how a supersonic jet behaves with time as it deflects from an obstacle situated at various distances from the axis of the jet. The laboratory analog of this phenomenon scales very well to the astrophysical case of a stellar jet which deflects from a molecular cloud core. An important component to our study was to expand the observational database of best astrophysical example, HH~110, by obtaining new spatially-resolved high-spectral resolution observations capable of distinguishing thermal motions from turbulent motions, and by acquiring a new deep infrared H$_2$ image that can be compared with existing optical emission line images from the Hubble Space Telescope.  The laboratory experiments span a range of times, spatial offsets between the axis of the jet and the center of the ball (impact parameters), viewing angles, opacities (backlighters), and materials. The experiments are reproducible and do not depend on composition or structure of foam, or on the pinhole diameter of the backlighter (spatial resolution of the experiment). Synthetic radiographs of the experiment from RAGE match the experimental data extremely well, both qualitatively as images and quantiatively with Fourier analysis. In fact, the agreement is so good that we have used the synthetic velocity maps from RAGE to compare the internal dynamics of the experiment with those that we measure from the new spectral maps of HH 110.  A new wide-field H$_2$ image supports a scenario where HH~110 represents the shocked `spray' that results from a glancing collision of the HH~270 jet with a molecular cloud core.  The H$_2$ in HH~110 is offset from the H$\\alpha$ toward the side closest to the molecular core, consistent with the deflected jet model. The H$_2$ images also uncovered two sources within the core that drive collimated jets, one a bright near-infrared source, and the other a highly-obscured source that appears to be a dense protostellar disk observed nearly edge-on, as in HH~30.  The experiments provide several important insights into how deflected supersonic jets like HH~110 behave. In the experiment, entrainment of material in the obstacle occurs in part because the morphology of the contact discontinuity between the shocked jet and shocked obstacle easily develops a complex 3-D structure of cavities that enables the jet to isolate clumps of obstacle material and entrain them into the flow. A similar process likely operates in HH~110. The experiments also reveal filamentary structure in the working surface area of the deflected bow shock, but the relative motion between these filaments is subsonic. Hence, while this dynamical process will generate density fluctuations in the outflowing gas, it cannot produce the filamentary structure and $\\sim$ Mach 5 shocks shown by the new velocity maps of HH~110. For this reason the best model for HH~110 remains that of a pulsed jet which interacts with a molecular cloud core.  Synthetic position-velocity maps along the deflected jet from the RAGE simulations of the experiments appear as a series of arcs, similar to those observed in astronomical observations of resolved bow shocks. A close examination of the experimental data shows that these arcs correspond to regions where the slit crosses different regions of the flow, such as cavities evacuated by the jet, the jet itself, or entrained material from the ball. This correspondance between the appearance of a p-v diagram and the actual morphology of a complex flow is an intuitive, although perhaps unexpected result of studying the dynamics within the experimental flow.  Finally, observations of the deflected bow shock from different viewing angles emphasize that the observed morphology and collimation properties of bow shocks depend strongly upon the orientation of the flow with respect to the observer. As one would expect, a bow shock deflected toward the observer appears less collimated than one that is redirected into the plane of the sky. The impact parameter of the jet and obstacle determines how much the jet deflects from the obstacle and how rapidly the obstacle becomes disrupted by the jet.  While experimental analogs of astrophysical jets are highly unlikely to ever reproduce accurate emission line maps, laser experiments can provide valuable insights into how the dynamics of complex flows behave. Our study of a deflected supersonic jet is only one example of how the fields of astrophysics, numerical computation, and laboratory laser experiments can compliment one another. We are currently embarking on a similar program to study the dynamics within supersonic flows that are highly clumpy, and other investigations are underway related to the launching and collimation of jets \\citep{bellan09}."
    },
    "0910/0910.5227_arXiv.txt": {
        "abstract": "We have studied the effects of electron--ion non-equipartition in the outer regions of relaxed clusters for a wide range of masses in the $\\Lambda$CDM cosmology using one-dimensional hydrodynamic simulations. The effects of the non-adiabatic electron heating efficiency, $\\beta$, on the degree of non-equipartition are also studied. Using the gas fraction $f_{\\rm gas} = 0.17$ (which is the upper limit for a cluster), we give a conservative lower limit of the non-equipartition effect on clusters. We have shown that for a cluster with a mass of $M_{\\rm vir}\\sim 1.2 \\times 10^{15} M_{\\odot}$, electron and ion temperatures differ by less than a percent within the virial radius $R_{\\rm vir}$. The difference is $\\approx 20\\%$ for a non-adiabatic electron heating efficiency of $\\beta \\sim 1/1800$ to $0.5$ at $\\sim 1.4 R_{\\rm vir}$. Beyond that radius, the non-equipartition effect depends rather strongly on $\\beta$, and such a strong dependence at the shock radius can be used to distinguish shock heating models or constrain the shock heating efficiency of electrons. With our simulations, we have also studied systematically the signatures of non-equipartition on X-ray and Sunyaev--Zel'dovich (SZ) observables. We have calculated the effect of non-equipartition on the projected temperature and X-ray surface brightness profiles using the MEKAL emission model. We found that the effect on the projected temperature profiles is larger than that on the deprojected (or physical) temperature profiles. The non-equipartition effect can introduce a $\\sim 10\\%$ bias in the projected temperature at $R_{\\rm vir}$ for a wide range of $\\beta$. We also found that the effect of non-equipartition on the projected temperature profiles can be enhanced by increasing metallicity. In the low-energy band $\\lesssim 1$~keV, the non-equipartition model surface brightness can be higher than that of the equipartition model in the cluster outer regions. Future X-ray observations extending to $\\sim R_{\\rm vir}$ or even close to the shock radius should be able to detect these non-equipartition signatures. For a given cluster, the difference between the SZ temperature decrements for the equipartition and the non-equipartition models, $\\delta\\Delta T_{\\rm SZE}$, is larger at a higher redshift.  For the most massive clusters at $z \\approx 2$, the differences can be $\\delta\\Delta T_{\\rm SZE} \\approx $ 4--5$~\\mu$K near the shock radius. We also found that for our model in the $\\Lambda$CDM universe, the integrated SZ bias, $ Y_{\\rm non{\\text -}eq} /Y_{\\rm eq}$, evolves slightly (at a percentage level) with redshift, which is in contrast to the self-similar model in the Einstein--de Sitter universe. This may introduce biases in cosmological studies using the $f_{\\rm gas}$ technique. We discussed briefly whether the equipartition and non-equipartition models near the shock region can be distinguished by future radio observations with, for example, the Atacama Large Millimeter Array. ",
        "introduction": "\\label{sec:intro} Observational and theoretical studies have shown that the study of the intracluster medium (ICM) can be used as a test of plasma physics under extreme environments that cannot be achieved in terrestrial laboratories, as well as an important cosmological probe.  If we assume the matter content of clusters is a fair sample of the universe, the baryon fraction of clusters can be used as an estimator of the average value for the universe, with proper correction for the baryons contained in the stellar component and for a small amount of baryonic matter expelled from clusters during the formation process.  Recent work from \\citep{ARS08} has shown that the combined results of the baryon fraction from X-ray observation with the cosmic microwave background (CMB) data can give powerful constraints on cosmological parameters, such as the equation of state of the dark energy.  However, the study of cosmology using clusters of galaxies relies heavily on the understanding of cluster physics.  For precision cosmology, systematic uncertainties at even the percent level are significant. For example, in current studies, the baryon fraction within clusters is assumed to be independent of redshift and the mass of clusters.  It would be important to see if these assumptions are justified; if not, it is important to study the dependence of the baryon fraction on cluster properties and redshift.  Even if the dependence on redshift is weak, the correction factor for the baryon content within clusters compared to the average value in the universe could affect the constraints on cosmological parameters.   Studying cluster outskirts ($\\gtrsim R_{200}$\\footnote{ $R_{\\Delta}$ is the radius within which the mean total mass density of the cluster is $\\Delta$ times the critical density. The virial radius $R_{\\rm vir}$ is defined as a radius within which the cluster is virialized.  For the Einstein--de Sitter universe, $ R_{\\rm vir} \\approx R_{178}$, while for the standard $\\Lambda$CDM universe, $ R_{\\rm vir} \\approx R_{95}$. }) is very important because the boundary conditions of the cluster outskirts constrain the global properties of a cluster \\citep[e.g.,][]{TSN00}. Also, the outer envelopes of clusters have been thought to be less subject to some additional physics including active galactic nucleus (AGN) feedback, and that the outer regions of clusters may provide better cosmological probes. Currently, there are very few observations of the properties of the ICM in the outer parts of clusters.  Thus, most of our understanding of these regions is still based on numerical hydrodynamic simulations which assume the hot plasma is a fluid.  In these simulations, the clusters are formed from mergers and accretion of dark matter and baryonic gas in overdense regions.  A variety of shocks with different geometries along the large-scale structure (LSS) filaments and transverse to them near and beyond the virial radius are unambiguous predictions of the cosmological hydrodynamic simulations.  Unfortunately, the lack of observational information on the clusters outskirts prevents us from understanding the accretion shock region, and hence the input physics for the numerical simulations is called into question.  For example, the thermodynamic state of the shocked gas, as well as the shock position, depends on the pre-shock gas temperature; the shock will be weaker if the infalling gas is pre-heated \\citep{TSN00}.   Even worse, recently it was noted that the non-fluid properties may be important in regions near the virial radius, where the Coulomb collisional mean free path is comparable to the cluster size of a few Mpc \\citep{Loe07}.  The Coulomb collisional timescale can also be of the order of the age of the cluster.  This suggests that a full kinetic gas theory is needed instead of the fluid approximation when studying the gas properties near the edge of the cluster.  Direct consequences include non-equipartition between electrons and ions \\citep{FL97, EF98}, element sedimentation \\citep{CL04}, and suprathermal evaporation of hot gas from the clusters (\\citealt{Loe07}; see also \\citealt{Med07}).  Some of these effects can lead to a bias in baryon fraction measurements, and hence cosmological studies.  Recently, progress has been made in the study the baryon content of the outer regions of clusters through the X-ray observations together with the Sunyaev--Zel'dovich (SZ) effect on the CMB by the hot electrons in the ICM out to $\\sim R_{200}$ \\citep{ALN+07}.  A 3$\\sigma$ result from the {\\it WMAP} three-year data suggests that $35\\%\\pm8\\%$ of the thermal energy in ICM are missing, indicating that the baryons in clusters may be missing even accounting for those locked in stars.  The result is also supported by independent measurements from other X-ray and SZ observations \\citep{Ett03, LBC+06, VKF+06, Evr+08}.  Using the X-ray observations together with numerical simulations, \\citet{Evr+08} reported that as much as $50\\%$ of the thermal energy can be missing in the ICM. Although \\citet{Gio+09} reported that the total baryon fraction within a smaller radius of $R_{500}$ of massive clusters are consistent to the cosmic value within $1 \\sigma$ when all the X-ray hot gas, stellar mass in galaxies, gas depletion during cluster formation, and intracluster light from stars are taken into account, if the missing baryons measured in the outer region ($\\lesssim R_{200}$) is really significant, this may indicate either a yet-unknown baryonic component, or some new astrophysical processes in the ICM which is driving out the gas from the clusters.  While there is no evidence for any undetected baryonic component, \\citet{ALN+07} pointed out that the missing of hot baryons can either be explained by the thermal diffusion or the evaporation of baryons out of the virial radius of clusters.  Another possibility is that electron temperature is lower than that of the equipartition value \\citep{WSL+08}.  Given the advancements in the X-ray and SZ observations of the cluster outer regions \\citep{ALN+07, Bau+09, GFS+09, Rei+09}, as well as the growing evidence of missing thermal energy in the ICM and the possible negative implications for cosmological tests, a more detailed study of the kinetic processes in cluster envelopes is necessary.  While magnetic fields may affect some of the kinetic effects of transport processes such as thermal conduction, the magnetic effects on non-equipartition should not be important since the physics is local.  Moreover, it is known that various astrophysical shocks in magnetized environment lead to non-equipartition \\citep{GLR07, HSL+01}.  The collisionless accretion shock at the outer boundary of a cluster should primarily heat the ions since they carry most of the kinetic energy of the infalling gas.  Assuming that cluster accretion shocks are similar to those in supernova remnants, the electron temperature, $T_e$, immediately behind the shock would be lower than the ion temperature, $T_i$.  The equilibration between electrons and ions would then proceed by Coulomb collisions.  Near the virial radius, due to the low density, the Coulomb collisional timescale can be comparable to the age of the cluster, and the electrons and ions may not achieve equipartition in these regions \\citep{FL97}.  Since X-ray and SZ observations measure the properties of the electrons in the ICM, the net effect is to underestimate the total thermal energy content within clusters.  This might account for some or all of the missing thermal energy in the ICM derived by the X-ray and SZ observations.  As mentioned above, non-equipartition of ions and electrons is observed in various astrophysical shocks.  Most supernova remnants with high Mach numbers comparable to cluster accretion shocks have electron temperatures which are lower than the ion temperatures \\citep{GLR07}; in situ measurements from satellites show the same feature in the Earth's bow shock \\citep{HSL+01}.  On the other hand, X-ray observations of the merger shock in the Bullet Cluster indicate that the equilibration time may be shorter than that expected from Coulomb collisions alone \\citep{MV07}.  However, this merger shock has a Mach number of a few.  From both supernova remnant measurements and a physical model, \\citet{GLR07} have shown that electron heating efficiency within a shock front (usually tens of the gyroradius) is inversely proportional to the Mach number squared.  If the results can also be applied to the ICM, the low Mach number merger shocks would be immediately heated to $T_e/T_i \\sim 1$, while cosmological accretion shocks with much higher Mach numbers would only be heated to $T_e/T_i \\ll 1$ by collisionless processes.  After the electrons and ions pass through the thin shock front, they will likely be equilibrated by Coulomb collisions alone \\citep{BDD08, BPP08}.  The non-equipartition in cluster of galaxies has been previously studied by \\citet{FL97} and \\citet{EF98} in semianalytic models.  They have shown that the temperature difference can be significant in the outer one-third of the shock radius of a cluster.  One- or three-dimensional simulations for some individual clusters have also been studied \\citep{CAT98, Tak99, RN09}.  While these simulations use different cluster or cosmological models, a general agreement is that the effect of non-equipartition is important if shock heating efficient of electrons is low ($\\ll 1$) and the equilibration afterward is due to Coulomb collisions alone.  In this paper, we study systematically the effects of non-equipartition on X-ray and SZ observables in outer regions of relaxed galaxy clusters, which is particularly important for cosmological studies.  We carry out one-dimensional hydrodynamic simulations with realistic Navarro-Frenk-White (NFW) model under the concordance $\\Lambda$CDM cosmological background to provide a sample of clusters (groups) with different masses ($10^{13}$--$10^{16} M_{\\odot}$) at different redshifts ($z=0$--$2$). Even though we are studying the kinetic non-fluid properties in the cluster outer regions, the hydrodynamic treatment in modeling the cluster dynamical properties is reasonable and is justified as follows. Even dynamically unimportant magnetic fields should be able to reduce significantly the diffusion mean free path perpendicular to the magnetic field \\citep{Sar86,BK09}. The suppression of diffusion in a plasma depends on the topology of magnetic fields.  For uniform magnetic fields, only diffusion perpendicular to the local magnetic field is suppressed, and along the field, particles move freely; their mean free path along a field line is still determined by Coulomb collisions. On large scales, diffusion is suppressed in bulk only if the magnetic fields are random and highly tangled on small scales. To include anisotropic diffusion in the calculation would be difficult since the magnetic field structure is not known well enough.  However, there is some evidence from large-scale magnetohydrodynamic simulations that magnetic fields in galaxy clusters are chaotic with correlation and reversal length scales of  $\\sim 50$ and $\\sim 100$~kpc, respectively \\citep{DBL02}. Hence, we simply assume that diffusion is suppressed.  We also assume that electrons and ions are equilibrated locally on a long Coulomb collisional timescale, and assume that equilibration via plasma instabilities \\citep{SCK+05, SCK+08} does not occur except at the shocks (see Section~\\ref{sec:e-heating}). Previous studies show that the dynamical properties of cluster outer regions in one-dimensional simulations successfully reproduce those simulated in three-dimensional calculations \\citep{NFW95, RK97}.  The advantages of the one-dimensional simulations for our problem are presented in Section~\\ref{sec:hydro}. We emphasize the signatures of non-equipartition on X-ray and SZ observations in our studies.  We also study the effect of electron shock heating efficiency on the degree of non-equipartition.  Thus, observations of electron--ion equilibration may give constraints to the electron heating efficiency, and hence the electron heating mechanism.  The wider parameter space compared to previous work explored in this paper allows us to study the impact of non-equipartition effects on cosmological studies in a future paper.  The paper is organized as follows.  In Section~\\ref{sec:hydro}, we describe the set up of our hydrodynamic models.  The detailed implementations of the shock heating and the Coulomb equilibration process for our simulations are presented in Section~\\ref{sec:e-heating}.  The ability of our simulations to reproduce analytic test models relevant to our studies is discussed in Section~\\ref{sec:test}.  We present the simulated dynamics of our realistic NFW cluster models in the standard $\\Lambda$CDM cosmology in Section~\\ref{sec:NFW-DE_dynamics}.  These cluster models are used to study the non-equipartition effects presented in the paper.  We define the X-ray and SZ observables for our models to be studied, and also present the results for these observables in Section~\\ref{sec:obs}.  We discuss and conclude our work in Section~\\ref{sec:conclusion}.  Unless otherwise specified, we assume the Hubble constant $H_0 = 71.9~h_{71.9}$~km~s$^{-1}$~Mpc$^{-1}$ with $h_{71.9}=1$, the total matter density parameter $\\Omega_{M,0} = 0.258$, the dark energy density parameter $\\Omega_{\\Lambda} = 0.742$, and the gas fraction $f_{\\rm gas}=\\Omega_b/\\Omega_M=0.17$, where $\\Omega_b$ is the baryon density parameter, for the realistic NFW model in the standard $\\Lambda$CDM cosmology\\footnote{\\tiny http://lambda.gsfc.nasa.gov/product/map/dr3/parameters\\_summary.cfm}, and a hydrogen mass fraction $X=76\\%$ for the ICM throughout the paper. ",
        "conclusions": "\\label{sec:conclusion} Using one-dimensional hydrodynamic simulations, we have calculated a sample of realistic NFW clusters in a range of masses in the $\\Lambda$CDM cosmology.  The cluster properties we simulated are consistent with the one-dimensional $N$-body simulations by \\citet{RK97}, and they have shown that their calculations reproduce the density and temperature profiles of the three-dimensional simulated relaxed clusters in the outer regions.  Our one-dimensional hydrodynamic simulations help us to isolate the important physical processes under controlled conditions. We have studied in detail the effect of non-equipartition in the outer regions of relaxed clusters in the $\\Lambda$CDM cosmology.  Using $f_{\\rm gas} = 0.17$ (which is the upper limit for a cluster), we give a conservative lower limit of the non-equipartition effect on clusters.  We have shown that for a cluster with a mass of $M_{\\rm sh}\\sim 1.5 \\times 10^{15} M_{\\odot}$, within $R_{\\rm vir}$, electron and ion temperatures only differ by less than a percent.  Our results show that the effect is smaller than those calculated from recent three-dimensional simulations, which shows that $T_e$ can be biased low by 5\\% at $R_{200} \\sim 0.7 R_{\\rm vir}$ \\citep[model CL104 in][]{RN09}.  A detailed analysis is needed to address the difference, but a possible explanation may be that in a three-dimensional cluster the accretion shock can be formed further in. Our results show that $T_e/{\\bar T}$ can reach $\\approx 0.8$ for a range of non-adiabatic electron heating efficiency $\\beta \\sim 1/1800$ to $0.5$ at $\\sim 0.9 R_{\\rm sh}$ (or $\\sim 1.4 R_{\\rm vir}$).  Beyond that radius, $T_e/{\\bar T}$ depends rather strongly on $\\beta$, and such a strong dependence at the shock radius can be used to distinguish shock heating models or constraint the shock heating efficiency of electron. We also show that the effect of non-equipartition is larger for more massive clusters, which is consistent to analytic self-similar models in the Einstein--de Sitter universe \\citep{FL97}.   Using the algorithm developed by \\citet{Vik06} which takes into account the soft emission at low temperature down to $\\sim 0.5$~keV, arbitrary metallicity, and instrumentation response, we calculated the X-ray spectroscopic-like temperature profiles which are the one to be directly determined observationally.  The effect of non-equipartition on the projected temperature profiles is larger than that on the deprojected (or physical) temperature profiles. Non-equipartition effects can introduce a $\\sim 10\\%$ bias in the projected temperature at $R_{\\rm vir}$ for a wide range of $\\beta$. This is because for the projected temperature profile, electrons in the outer region also contribute to the inner region. We also found that the effect of non-equipartition on the projected temperature profiles can be enhanced by increasing metallicity. This is because the domination of the line emissions in the soft band spectra is enhanced by increased metallicity, which is more important for the non-equipartition model where the electron temperature is lower.    The effect of non-equipartition on X-ray surface brightness profile in the 0.3--2~keV band is smaller than that on the projected temperature profile.  This is because the surface brightness depends on density squared but with a weaker dependence on temperature.  This means that in the outer regions, clusters in non-equipartition have similar X-ray surface brightness profile for $E \\lesssim 2$~keV, but with bigger difference in temperature compared to those equipartition counterparts. For $E \\gtrsim 2$~keV, non-equipartition effects on X-ray surface brightness profiles are similar to those on the projected temperature profiles. For a cluster with $M_{\\rm sh}\\sim 1.5 \\times 10^{15} M_{\\odot}$, the effect of non-equipartition on surface brightness profiles in all energy bands is $\\lesssim 10\\%$ for radii $\\lesssim 3$~Mpc; beyond that, the effect can be important.  We found that for the non-equipartition model, the surface brightness profile in the low-energy band $\\lesssim 1$~keV can be higher than that of the equipartition model in the cluster outer regions.  Non-equilibrium ionization, which was not considered in our calculations of the emissivities, may further enhance the line emissions in the soft bands ($E \\lesssim 1$~keV).   Current X-ray observations extend to only $\\sim R_{200} \\sim 2$~Mpc, although some results from recent $Suzaku$ observations begin to go a bit beyond that \\citep{GFS+09}. Within those regions with $\\lesssim R_{200}$, electrons and ions should be almost in equipartition and the signatures in the X-ray temperature and surface brightness should be rather weak.  But future X-ray observations may extend to $\\sim R_{\\rm vir} \\approx 1.4 R_{200}$ or even close to the shock radius. We have shown that non-equipartition of electrons and ions should be detectable in those studies.  The results by \\citet{GFS+09} support our conclusion.   The effects of non-equipartition on the SZ effect were studied.   At $z=0$, the effect on the Comptonization parameters is similar to that of the projected temperature profiles.  For a cluster with $M_{\\rm sh}\\sim 1.5 \\times 10^{15} M_{\\odot}$, $y_{\\rm non{\\text -}eq}/y_{\\rm eq} \\approx 0.93 (0.8)$ at $1 (1.3) R_{\\rm vir}$.  For a given cluster, the difference between the SZ temperature decrements for the equipartition and the non-equipartition models is larger at a higher redshift.  For the most massive clusters at $z \\approx 2$,   the differences can be $\\delta\\Delta T_{\\rm SZE} \\approx$~4--5$~\\mu$K near the shock radius.  A detailed analysis of whether the equipartition and non-equipartition models near the shock region can be distinguished by, for example, ALMA, will be presented in a future paper.    The effects on the integrated SZ Comptonization parameter, which measures the thermal energy content of the electrons, were studied. We have shown that the integrated SZ bias, $Y_{\\rm non{\\text -}eq}/Y_{\\rm eq}$, increases as the cluster mass increases, which is expected as the effect of non-equipartition increases with mass. In general, the non-equipartition effect is larger for the realistic NFW model in the $\\Lambda$CDM universe than that for the self-similar model in the Einstein--de Sitter universe, assuming that they have the same $f_{\\rm gas}$. Our simulations suggest that for relaxed clusters with $M_{\\rm sh}\\sim 1.5 \\times 10^{15} M_{\\odot}$, the non-equipartition effect can account for only about 2\\%--3\\% of the missing thermal energy globally. For the most massive clusters, up to 4\\%--5\\% of the thermal energy beyond the equipartition value may be stored in the thermal energy of ions near the shock radius, but for clusters with $M_{\\rm sh} \\lesssim 5 \\times 10^{14} M_{\\odot}$, the non-equipartition effect is less than $1 \\%$. Thus, we argue that, at least for relaxed clusters, the non-equipartition effect alone can only account for some of the missing thermal energy problem, if any, for high-mass clusters but not for clusters with smaller masses.  On the other hand, this suggests that hot gas may be missing due to other astrophysical processes not yet known, and the $f_{\\rm gas}$ in a real cluster should be lower than that we used in our numerical simulations. We have estimated that reducing $f_{\\rm gas}$ by $20\\%$ will enhance the local non-equipartition effect near the outer region by a few percent, but the integrated SZ bias, $Y_{\\rm non{\\text -}eq}/Y_{\\rm eq}$ is not affected by more than a percent.  We emphasis here that even though non-equipartition effects may not affect the global energy budget significantly, the effect is still important locally in the outer regions ($\\sim R_{\\rm vir}$) of a cluster.  Future X-ray and SZ observations may extend out to $R_{\\rm sh}$, and the effect of non-equipartition should be considered when studying cluster properties in those regions.   We found that for our realistic NFW model in the $\\Lambda$CDM universe, $Y_{\\rm non{\\text -}eq}/Y_{\\rm eq}$ evolves with redshift, which is in contrast to the self-similar model in the Einstein--de Sitter universe.  For our realistic NFW model in the $\\Lambda$CDM universe, $Y_{\\rm non{\\text -}eq}/Y_{\\rm eq}$ decreases as $z$ decreases. This is probably due to the decreasing rate of accretion onto clusters in the $\\Lambda$CDM universe during the period of cosmological acceleration, which results in a relatively longer timescale for the electron--ion equilibration inside a cluster compared to a cluster with the same mass in the Einstein--de Sitter universe. Though the magnitude of $Y_{\\rm non{\\text -}eq}/Y_{\\rm eq}$ is small for the range of cluster masses, even a percentage level deviation in the most massive clusters can be important for precision cosmology studies.  Such a variation of $Y$ with $z$ would introduce an apparent evolution in $f_{\\rm gas}$, which would bias the cosmological studies using the $f_{\\rm gas}$ techniques \\citep{ARS08}. Recently, \\citet{RN09} have shown that the non-equipartition effect on $Y$ can be enhanced by major mergers up to $30\\%$, although for low Mach number mergers, the shock heating efficiency for electrons may be higher which can weaken the non-equipartition effect \\citep{GLR07, MV07}.  The temporary boost due to mergers may have a significant effect on the estimation of cosmological parameters using clusters. We defer a detailed study of the effect on cosmology studies in a future paper.    K.W. thanks Avi Loeb and Brian Mason for helpful discussions. Support for this work was provided by the National Aeronautics and Space Administration, through {\\it Chandra} Award Numbers TM7-8010X, GO7-8135X, GO8-9083X, and GO9-0035X, NASA $XMM-Newton$ grants NNX08AZ34G, NNX08AW83G, and NASA $Suzaku$ grant NNX08AI27G. We thank the referee for helpful comments.   \\appendix"
    },
    "0910/0910.0258_arXiv.txt": {
        "abstract": "We describe and test a new method for creating initial conditions for cosmological N-body dark matter simulations based on second-order Lagrangian perturbation theory (2lpt). The method can be applied to multi-mass particle distributions making it suitable for creating resimulation, or `zoom' initial conditions. By testing against an analytic solution we show that the method works well for a spherically symmetric perturbation with radial features ranging over more than three orders of magnitude in linear scale and eleven orders of magnitude in particle mass. We apply the method and repeat resimulations of the rapid formation of a high mass halo at redshift $\\sim50$ and the formation of a Milky-Way mass dark matter halo at redshift zero.  In both cases we find that many properties of the final halo show a much smaller sensitivity to the start redshift with the 2lpt initial conditions, than simulations started from Zel'dovich initial conditions.  For spherical overdense regions structure formation is erroneously delayed in simulations starting from Zel'dovich initial conditions, and we demonstrate for the case of the formation the high redshift halo that this delay can be accounted for using the spherical collapse model.  In addition to being more accurate, 2lpt initial conditions allow simulations to start later, saving computer time. ",
        "introduction": "Computer simulations of cosmological structure formation have been crucial to understanding structure formation particularly in the non-linear regime. Early simulations e.g. \\cite{Aarseth79} used only around a thousand particles to model a large region of the Universe, while recent simulations of structure formation have modelled over a billion particles in just a single virialised object \\citep{Springel08,Stadel09}.  Since the advent of the Cold Dark Matter (CDM) model \\citep{Peebles82,DEFW85}, most computational effort has been expended modelling structure formation in CDM universes. Early work in the 1980s focused on `cosmological' simulations where a representative region of the universe is modelled, suitable for studying large-scale structure \\citep{DEFW85}. The starting point for these CDM simulations, the initial conditions, are Gaussian random fields.  The numerical techniques needed to create the initial conditions for cosmological simulations were developed in the 1980s and are described in \\cite{Efstathiou85}.  As the algorithms for N-body simulations have improved, and the speed of computers increased exponentially with time, it became feasible in the 1990s to simulate the formation of single virialised halos in the CDM model with enough particles to be able to probe their internal structure \\citep[e.g.][]{NFW96, NFW97, Ghigna98, Moore99}. These more focused simulations required new methods for generating the initial conditions.  The algorithm of \\cite{HoffmanRibak91} for setting up constrained Gaussian random fields was one method which could be applied to setting up initial conditions for a rare peak where a halo would be expected to form.  The essence of this technique is that it allows selection of a region based on the properties of the linear density field.  However, the objects we actually observe in the Universe are the end products of non-linear structure formation and it is desirable, if we want to understand how the structure we see formed, to be able to select objects on the basis of their final properties.  This requirement led to an alternative method for producing initial conditions based first on selecting objects from a completed simulation e.g. \\cite{Katz93}, \\cite{Navarro94}.  The initial conditions for the first or parent simulation, were created using the methods outlined by \\cite{Efstathiou85}. The density field is created out of a superposition of plane waves with random phases.  Having selected an object at the final (or any intermediate) time from the parent simulation a fresh set of initial conditions with higher numerical resolution in the region of interest, which we will call `resimulation' initial conditions (also called `zoom' initial conditions) can be made. Particles of different masses are laid down to approximate a uniform mass distribution, with the smallest mass particles being concentrated in the region from which the object forms. The new initial conditions are made by recreating the the same plane waves as were present in parent simulation together with the addition of new shorter wavelength power.  An alternative technique for creating resimulation initial conditions was devised by \\cite{Bertschinger01} based on earlier work by \\cite{Salmon96} and \\cite{Pen97} where a Gaussian random field with a particular power spectrum is created starting from a white noise field.  Recently parallel code versions using this method has been developed to generate very large initial conditions \\citep{Prunet08,Stadel09}. A feature common to both these techniques, to date, is that the displacements and velocities are set using the Zel'dovich approximation \\citep{Zeldovich70}, where the displacements scale linearly with the linear growth factor.  It has long been known that simulations starting from Zel'dovich initial conditions exhibit transients and that care must be taken choosing a sufficiently high start redshift so that these transients can decay to a negligible amplitude \\citep[e.g][]{Efstathiou85}. A study of the behaviour of transients using Lagrangian perturbation theory by \\cite{Scoccimarro98}, showed that for initial conditions based on second-order Lagrangian perturbation theory, the transients are both smaller and decay more rapidly than first-order, or Zel'dovich, initial conditions.  \\cite{Scoccimarro98}  gave a practical method for implementing second-order Lagrangian initial conditions. The method has been implemented in codes for creating cosmological initial conditions by \\cite{Sirko05} and \\cite{Crocce_etal06}.  These codes are suitable for making initial conditions for cosmological simulations, targeted at large-scale structure, but they do not allow the creation of resimulation initial conditions, the focus of much current work on structure formation.  In this paper we describe a new method for creating second-order Lagrangian initial conditions which can be used to make resimulation initial conditions.  The paper is organised as follows: in Section~\\ref{SECTMOT} we introduce 2lpt theory initial conditions and motivate their advantages for studying structure in the non-linear regime by applying them to the spherical top-hat model; in Section~\\ref{MAKERESIMS} we describe how Zel'dovich resimulation initial conditions are made, and the new method for creating 2lpt initial conditions, and test the method against an analytic solution for a spherically symmetric perturbation; in Section~\\ref{REALAPPLICATIONS} we apply the method and analyse the formation of a dark matter halo at high and low redshift for varying starting redshifts for both Zel'dovich and 2lpt initial conditions; in Section~\\ref{SECTSUMMARY} we summarise and discuss the main results; in the Appendix we evaluate two alternate interpolation schemes which have been used in the process of making resimulation initial conditions. ",
        "conclusions": "}  We have implemented a new method to make second order Lagrangian perturbation theory (2lpt) initial conditions which is well suited to creating resimulation initial conditions. We have tested the method for an analytic spherically symmetric perturbation with features ranging over a factor of 5000 in linear scale or eleven orders of magnitude in mass, and can reproduce the analytic solution to a fractional accuracy of better than one percent over this range of scales in radius measured from the symmetry centre.  Applying the new method we have recreated the initial conditions for the formation of a high redshift halo from \\cite{Gao05} and the a Milky-way mass halo from \\cite{Springel08}.  We find from studying the properties of the final halos, that the final conditions show much less sensitivity to the start redshift when using 2lpt initial conditions than when started from Zel'dovich initial conditions.  For didactic purposes, we have calculated the effect of using Zel'dovich and 2lpt initial conditions for a spherical top-hat collapse. In this simple case, the epoch of collapse is delayed by a timing error which depends primarily on the number of expansions between the epoch of the initial conditions and the collapse time. This timing error, for fixed starting redshift, grows with collapse time for Zel'dovich initial conditions, but decreases with collapse time for 2lpt initial conditions. Applying the top-hat timing errors as a correction to the radius of shells containing fixed amounts of mass as a function of time has the effect of bringing the Zel'dovich and 2lpt resimulations of the high redshift \\cite{Gao05} halo into near coincidence.  We find that for the Aq-A-5 halo the choice of start redshift for the Zel'dovich initial conditions in \\cite{Springel08} was sufficiently high at $z_s=127$ to make little difference to the properties of the halo at $z=0$ when compared to runs starting from 2lpt initial conditions.  For lower starting redshifts ($z_s = 63, 31, 15$) the positions of substructures at $z=0$ do become sensitive to the precise start redshift.   \\cite{Knebe09} have also looked at the properties of halos, for the mass range ($10^{10} - 10^{13}h^{-1}M_\\odot$) at redshift zero, for simulations starting from Zel'dovich and 2lpt initial conditions (generated using the code by \\cite{Crocce_etal06}).  In their paper they simulated a cosmological volume which yielded many small halos and about ten halos with 150000 or more particles within the virial radius. In their study they looked at the distribution of halo properties including the spin, triaxiality and concentration.  They also matched halos in a similar way to the way we have matched subhalos and looked at the ratios of masses, triaxialities, spins and concentrations of the matching halos. The authors concluded that any actual trends with start redshift or type of initial condition by redshift zero are certainly small (for start redshifts of 25 - 150). It is not possible to directly compare our results with theirs, but the trends in the halo properties we have observed in halo Aq-A-5 at redshift zero are indeed small, and not obviously in disagreement with their results for populations of halos.  For many purposes, the differences which arise from using Zel'dovich or 2lpt initial conditions may be small when compared to the uncertainties which arise in the modelling of more complex physical processes.  From this point of view the main advantage of using 2lpt initial conditions is that they allow the simulations to start later thus saving computer time. The issue of exactly what redshift one can start depends on the scientific problem, but the spherical collapse model, discussed in Subsection~\\ref{substh} gives a quantitative way to compare the start redshift for Zel'dovich or 2lpt initial conditions.  However it remains the case that it is advisable to test the sensitivity of simulation results to the start redshift by direct simulation as for example in \\cite{Power03} for resimulations of galactic dark matter halos."
    },
    "0910/0910.4198_arXiv.txt": {
        "abstract": "We report on the Camera Materials Test Chamber, a multi-vessel apparatus which analyzes the outgassing consequences of candidate materials for use in the vacuum cryostat of a new telescope camera.  The system measures the outgassing products and rates of samples of materials at different temperatures, and collects films of outgassing products to measure the effects on light transmission in six optical bands.  The design of the apparatus minimizes potential measurement errors introduced by background contamination. ",
        "introduction": "Introduction}  The Large Synoptic Survey Telescope (LSST) is a planned facility which will repeatedly and deeply image large (10 square degree) sections of the night sky at optical and very near infrared wavelengths to an unprecidented depth, gathering important data to address outstanding questions in cosmolgy and astrophysics\\cite{LSSTover}.  The etendue, or light gathering ability, of LSST, as measured by the primary mirror aperture size multiplied by the field of view, is more than an order of magnitude higher than any existing or planned optical facility.  Such imaging requires a 3.2 billion pixel cryogenic CCD camera maintained in a clean vacuum environment.  The cryostat of the LSST camera will measure 2.9 $m^{3}$, contain nearly 1000 kg of material, and have more than 40 $m^{2}$ of exposed surface area within, including the cold focal plane CCD detectors, two stages of supporting electronics, mechanical and thermal control structures, electrical cabling, electrical and fluid connectors and feed-thrus, and devices for internal metrology.  All materials within the evacuated cryostat will remain there for years and be subject to thermal cycling.  These materials must not outgass in a way that compromises the insulating vacuum nor change the net light transmittance to the focal plane between calibrations.  There is scant data in the literature in regard to how the outgassing and deposition of commercial and other materials may affect light transmittance to optical surfaces in vacuum.  Abromovici et al. limited the effect of elastomer outgassing on mirror reflectivities\\cite{ocon}.  They use resonant cavities pumped by a laser, thus measuring at only one fequency, and do not cool the optical collecting surfaces (mirrors).  Others have measured the outgassing and deposition rates of various matrials for many applications.  However, given the almost endless variety of available materials, achieving a database relevant to very different applications is impossible.  We have designed, contstructed, and commissioned an apparatus, the Camera Materials Test Chamber (CMTC), to quantiy the suitability of candidate materials for inclusion in the LSST camera cryostat.  The CMTC could also be used to perform suitability tests for any similar vacuum application, and, given its ability to isolate functions and cold stages, can be used for a variety of precision vacuum measurements, such as the performance of getters.  We report on the design and performance of the instrument, highlighting vacuum and thermal engineering considerations that may be of use to experimentalists. ",
        "conclusions": ""
    },
    "0910/0910.4960_arXiv.txt": {
        "abstract": "For half a century, evidence has been growing that the formation of stars follows a universal distribution of stellar masses. In fact, no stellar population has been found showing a systematic deviation from the canonical initial mass function (IMF) found for example for the stars in the solar neighbourhood. The only exception may be the young stellar discs in the Galactic Centre, which have been argued to exhibit a top-heavy IMF.  Here we discuss the question whether the extreme circumstances in the centre of the Milky Way may be the reason for a significant variation of the IMF. By means of stellar evolution models using different codes, we show that the observed luminosity in the central parsec is too high to be explained by a long-standing top-heavy IMF as suggested by other authors, considering the limited amount of mass inferred from stellar kinematics in this region. In contrast, continuous star formation over the Galaxy's lifetime following a canonical IMF results in a mass-to-light ratio and a total mass of stellar black holes (SBHs) consistent with the observations. Furthermore, these SBHs migrate towards the centre due to dynamical friction, turning the cusp of visible stars into a core as observed in the Galactic Centre. For the first time here we explain the luminosity and dynamical mass of the central cluster and both the presence and extent of the observed core, since the number of SBHs expected from a canonical IMF is just enough to make up for the missing luminous mass.  We conclude that observations of the Galactic Centre are well consistent with continuous star formation following the canonical IMF and do not suggest a systematic variation as a result of the region's properties such as high density, metallicity, strong tidal field etc. If the young stellar discs prove to follow a top-heavy IMF, the circumstances that led to their formation must be very rare, since these have not affected most of the central cluster. ",
        "introduction": "In \\citeyear{s55}, Salpeter found that the initial mass distribution of field stars in the range $0.4\\simless M_{\\star}/\\msun \\simless 10$ is a power-law with exponent 2.35. Since then, a large number of publications have investigated the initial mass function (IMF) of stars and made clear that star formation in general follows the same empirical law, the canonical IMF \\citep[and references therein]{k01}.  Due to its extreme conditions (mass density, velocity dispersion, tidal forces), the Galactic Centre provides a unique environment for testing the universality of the IMF. Star formation in the central region has thus been studied in detail, however no agreement has been reached on the nature of the IMF in either theory or observations: \\citet{mmt07} find a best fit of observations in the central parsec of our Galaxy with a model of constant star formation with a top-heavy IMF. On the other hand, \\citet{bse09} show that the old stellar cluster in the Galactic Centre very well resembles the bulge population. Observations of the young, massive Arches cluster in the central region of the Milky Way have long been interpreted as a prime example for top-heavy star formation \\citep[e.g.][]{fetal99,setal02,kfkn06,ksj07}. However, \\citet{esm09} have shown that a canonical IMF cannot be excluded for this cluster. \\citet{pau06} suggested a flat IMF for the young OB-stars observed in discs in the central parsec from the analysis of the K-band luminosity function. Based on more recent spectroscopic observations, \\citet{bmt+09} find strong evidence for this to be true. \\citet{br08} found from SPH simulations that the IMF of stars forming in fragmenting accretion discs strongly depends on the parameters of the underlying gas infall scenario. Unfortunately, theoretical IMF predictions have failed in the past to correctly describe the observations near the Galactic Centre \\citep{k08}.  In this paper, we combine observational data with models of stellar evolution and dynamics to constrain the stellar mass function and star formation history in the Galactic Centre. It is organised as follows: In Section~\\ref{sec:ssemodels}, we analyse the properties of models of the Galactic Centre assuming different star formation histories, and compare them to the observations. Section ~\\ref{sec:massprof} describes the mass profile of the central parsec and the effect of mass segregation. We discuss the IMF of the young stellar discs around \\sgra\\ and appropriate formation scenarios in Section~\\ref{sec:discimf} and summarise in Section~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  Various attempts have been made to study star formation in the Galactic Centre, but so far neither theory or simulations nor observations led to an agreement on the (initial) distribution of stellar masses. Here we have shown that theory and observations are consistent with star formation generally following a canonical IMF \\citep{k01} in the Galactic Centre, just as anywhere else in the Universe. Our main results can be summarised as follows: \\renewcommand{\\labelenumi}{\\arabic{enumi}.} \\begin{enumerate} \\item The mass-to-light ratio of the central parsec of the Milky Way is consistent with a constant or exponentially decreasing star formation rate following a canonical IMF. Models of constant star formation following an IMF with $\\alpha = 1.35$ are consistent with the observed luminosities but create $\\sim 10^5$ SBHs in the central parsec, ten times more than expected by other authors. Mass functions flatter than $\\alpha \\approx 1$ can be safely ruled out, since they cannot explain the observed number of bright stars and the diffuse light. \\item The core observed in the luminosity distribution with a radius of $r_{\\rm break} \\approx 0.2$\\,pc does not imply a core in the mass profile, but can be well explained by mass segregation as suggested by \\citet{fak06}, where dark remnants mass-dominate this region. Again, the observations are best explained by a canonical IMF, and are not compatible with star formation following a top-heavy IMF with $\\alpha \\le 1.35$, as this would create a core radius one order of magnitude larger. \\item Recent observations revealing a deficit of B-type stars in the young stellar discs suggest a top-heavy IMF for this population, which may be explained by tidal stripping of an infalling mass-segregated cluster, or unusual modes of star formation in a fragmenting accretion disc. However, these results do not allow conclusions on star formation in the Galactic Centre in general, for which we have no reason to assume it to be non-canonical. \\end{enumerate}  While other authors generally predicted a flat IMF \\citep[e.g.][]{ksj07} or a higher low-mass cut-off \\citep[e.g.][]{m93,lb03,l_06} from state-of-the-art theoretical star formation models of the Galactic Centre, we find that observations suggest star formation follows a canonical IMF even under the extreme circumstances present in the central cluster. This universality of the IMF poses a major challenge to our understanding of star formation processes (see also \\citealp{k01,k08b}).  Our results rely on the assumption that the stars observed within 1\\,pc from the Galactic Centre also formed there, suggesting star formation with a canonical IMF even under such exotic conditions. It is possible that some of the stars were brought in by massive star clusters which spiralled towards the SMBH through dynamical friction \\citep{pbmmhe06, fifm08}. However, this scenario is unlikely because a star cluster is stripped on its way towards the centre and loses mostly low-mass stars, since it would be in a mass segregated state soon after its formation. Therefore, the most likely scenario is a central cluster that formed over a Hubble time with a canonical IMF, where the very young stellar population of the stellar discs observed to have a very top-heavy IMF \\citep{bmt+09} constitutes a rare star formation event not typical for the bulk stellar population in the central cluster."
    },
    "0910/0910.4367_arXiv.txt": {
        "abstract": "We report the discovery of radio afterglow emission from the gamma-ray burst GRB\\,090423, which exploded at a redshift of 8.3, making it the object with the highest known redshift in the Universe. By combining our radio measurements with existing X-ray and infrared observations, we estimate the kinetic energy of the afterglow, the geometry of the outflow and the density of the circumburst medium. Our best fit model is a quasi-spherical, high-energy explosion in a low, constant-density medium. \\event\\ had a similar energy release to the other well-studied high redshift GRB 050904 ($z=6.26$), but their circumburst densities differ by two orders of magnitude. We compare the properties of \\event\\ with a sample of GRBs at moderate redshifts. We find that the high energy and afterglow properties of \\event\\ are not sufficiently different from other GRBs to suggest a different kind of progenitor, such as a Population III star.   However, we argue that it is not clear that the afterglow properties alone can provide convincing identification of Population III progenitors. We suggest that the millimeter and centimeter radio detections of \\event\\ at early times contained emission from a reverse shock component. This has important implications for the detection of high redshift GRBs by the next generation of radio facilities. ",
        "introduction": "\\label{sec:intro}  Because of their extreme luminosities GRBs are detectable out to large distances by current missions, and due to their connection to core collapse SNe \\citep{wb06}, they could in principal reveal the stars that form from the first dark matter halos ($z\\sim$ 20--30) through to the epoch of reionization at $z=11\\pm 3$ and closer \\citep{lr00,cl00b,gmaz04,ioc07}. As bright continuum sources, GRB afterglows  also make  ideal backlights to probe the intergalactic medium as well as the interstellar medium in their host galaxies.   Predicted to occur at redshifts beyond those where quasars are expected, they could be used to study both the reionization history and metal enrichment of the early universe \\citep{tkk+06}.  The fraction of detectable GRBs that lie at high redshift ($z>6$) is, however, expected to be small ($<$10\\%; \\citealt{pcb+09,bl06}). Until recently there were only two GRBs with measured redshifts $z>6$;  GRB\\,050904 \\citep{kka+06} and GRB\\,080913 \\citep{gkf+09} with $z=6.3$ and $z=6.7$, respectively. However, on April 23, 2009 the \\Swift\\ Burst Alert Telescope (BAT) discovered \\event\\ and the on-board X-ray Telescope (XRT) detected and localized a variable X-ray afterglow \\citep{tfl+09,sdc+09}. In ground-based follow-up observations no optical counterpart was found but a fading afterglow was detected by several groups at wavelengths longward of J band (1.2 $\\mu$m). Based on both broadband photometry and near infrared (NIR) spectroscopy, the sharp optical/NIR drop off was argued to be due to the Lyman-$\\alpha$ absorption in the intergalactic medium, consistent with a redshift with a best-fit value of $z=8.26^{+0.07}_{-0.08}$ \\citep{tfl+09}.  The high redshift of \\event\\ makes it the most distant observed GRB, as well as the most distant object of any kind other than the Cosmic Microwave Background.  This event occurred approximately 630 million years after the Big Bang, confirming that massive stellar formation occurred in the very early universe.  In this paper we report the discovery of the radio afterglow from \\event\\ with the Very Large Array\\footnote{The Very Large Array is operated by the National Radio Astronomy Observatory, a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.} (VLA).   Broadband afterglow observations provide constraints on the explosion energetics, geometry, and  immediate environs of the progenitor star. The afterglow has a predictable temporal and spectral evolution that depends on the kinetic energy and geometry of the shock, the density structure of the circumburst environment, and shock microphysical parameters which depend on the physics of particle acceleration and the circumburst magnetic field. To the degree that we can predict differences in the explosion and circumburst media between GRB progenitors at high and low redshifts, we can search for these different signatures in their afterglows. This has been the motivation for previous multi-wavelength modeling of the highest-$z$ afterglows \\citep{fck+06,gfm07}.  In order to investigate the nature of the \\event\\ explosion, we combine our radio measurements with published X-ray and NIR observations, and apply a model of the blast wave  evolution to fit the afterglow data (\\S\\ref{sec:results}).   We compare the explosion energetics, circumburst density, and other derived characteristics to a sample of well-studied events, and discuss prospects for using afterglow measurements to investigate the nature of high-$z$ massive star progenitors. ",
        "conclusions": "\\label{sec:dis}  \\event\\ is the highest-redshift object for which we have multi-wavelength observations, including good quality radio measurements. Below we address the following questions: based on its afterglow properties what can we learn about properties of the explosion and environs for this highest-redshift GRB? And, can we identify any differences between high and low redshift GRBs which indicate that they might arise from different progenitors? In particular, the initial generations of stars in the early universe are thought to be brighter, hotter and more massive ($>100\\, M_\\odot$) than stars today \\citep{hai08,byhm09}.   Detecting these so-called Population III (Pop III) stars is one of the central observational challenges in modern cosmology,  and the best prospect appears to be through observing their stellar death \\citep{hfw+03} via a supernovae (SNe) or gamma-ray burst explosion.   It is worth asking what observational signatures could signal a Pop III GRB.   Other than \\event ,  only one other $z > 6$ event, GRB\\,050904 ($z=6.26$), has high quality broadband afterglow measurements. In Fig.~\\ref{fig:comparison} we plot the best-fit parameters of these two GRBs along with a sample of well-studied lower redshift events from \\citet{pk01}. Both high redshift bursts stand out in terms of their large blast wave energy ($> 10^{52}$ erg). We know from samples of well-studied afterglows \\citep{fks+01,pk01,yhsf03}, that most  have radiative and kinetic energies of order $\\sim 10^{51}$ erg.  In the collapsar model the jet kinetic energy from a Pop III GRB could be 10--100 times larger than a Population II (Pop  II) event  \\citep{fwh01,hfw+03}.   However, an energetic explosion does not appear to be an exclusive property of high-$z$ GRBs. There is a small but growing population of bursts with energy $> 10^{52}$ erg, termed 'hyper-energetic GRBs' \\citep{cfh+09}, which includes moderate-$z$ events like GRB\\,070125 \\citep{ccf+08}  and GRB\\,050820A \\citep{ckh+06}.  Another potentially useful diagnostic is the density structure in the immediate environs of the progenitor star. The radio data is a sensitive {\\it in situ} probe of the density because its emission samples the optically thick part of the synchrotron spectrum. The afterglows of \\event\\ and GRB\\,050904 are best fit by a constant density medium and not one that is shaped by stellar mass loss \\citep{cl99}.  However, many afterglows at all redshifts are best fit by a constant density medium (e.g.~\\citealt{yhsf03}). The density obtained for GRB\\,050904 was the highest seen ($n\\approx84-680$ cm$^{-3}$) for any GRB to date, while \\event\\ with $n=0.9$ cm$^{-3}$ does not stand out (Fig.~\\ref{fig:comparison}), indicating these two high redshift bursts exploded in very different environments. A circumburst density of order unity is predicted for Pop III stars, since this density is  limited by strong radiation pressure in the mini halo from which the star was formed \\citep{byh03}.   This is not an unique property, since many local SNe explode in tenuous media, and so density constraints are not useful to signal Pop III explosions.  For the other afterglow parameters ($p$, $\\epsilon_e$, $\\epsilon_B$ and $\\theta_j$) there are no published predictions for how they may differ between different progenitor models. Thus we turn to considering the prompt high-energy emission of \\event.  \\citet{sdc+09} and \\citet{tfl+09} both noted that the high-energy properties of \\event\\ (fluence, luminosity, duration, radiative energy) are not substantially different from those moderate-$z$ GRBs.  We are not aware of any quantitative predictions based on these observed properties that would discriminate between Pop\\,II and Pop\\,III progenitors. For example, apart from the effect of ($1+z$) time dilation, there is no reason to expect high-$z$ GRBs to have significantly longer intrinsic durations. Collapsar models which form black holes promptly through accretion onto the proto-neutron star (Type I) or via direct massive ($>260$ M$_\\odot$) black hole formation (Type III) have durations set by jet propagation and disk viscosity timescales, respectively \\citep{mwh01,fwh01}, which are $\\sim 10$ s in the rest frame (with large uncertainties; \\citealt{fwh01}).  Metallicity can also be an important discriminant. There is a critical metallicity ($Z>10^{-3.5}\\,Z_\\odot$) below which high-mass Pop III stars dominate \\citep{byhm09,bl06}.  The contribution of Pop III stars to the co-moving star formation rate is expected to peak around $z=15$ but their redshift distribution exhibits a considerable spread to $z\\sim$7. Thus we might find high-redshift GRBs with Pop III progenitors in ``pockets'' of low metallicity. \\citet{sdc+09} argue for a lower bound of $Z>0.04\\,Z_\\odot$ based on their detection of excess soft X-ray absorption by metals along the line of sight, in comparison to the Milky Way column density predicted from \\ion{H}{1} (21\\,cm) measurements. We do not consider this a robust measurement as it is sensitive to a range of unaccounted-for systematic effects, including: spectral variability; spectral curvature; low-amplitude X-ray flares; and the presence of intervening (cosmological) absorption systems along the line of sight.  Summarizing the above discussion, we do not find that the individual properties of \\event\\ are sufficiently dissimilar to other GRBs to warrant identifying it as anything other than a normal GRB. We lack robust predictions of well-defined afterglow signatures that could allow us to unambiguously identify a Pop III progenitor star from its afterglow properties alone. Significantly larger numbers of GRBs at high redshift with  well-sampled afterglow light curves, high-resolution spectra, and host galaxy detections are needed to determine if high redshift GRB progenitors differ in a statistical sense from those at low redshift.  We note that, like GRB\\,050904, the \\event\\ afterglow indicates the signature of  reverse shock (RS) emission in the radio, as seen in the VLA  and  PdBI data. \\citep{ioc07} have studied the expected RS emission at high redshift, and they find that the effects of time dilation almost compensate for frequency redshift , resulting in a near-constant observed peak frequency in the mm band ($\\nu \\sim 200$~GHz) at a few hours post-event, and a flux at this frequency that is almost independent of redshift. Further, the mm band does not suffer significantly either from extinction (in contrast to the optical) or scintillation (in contrast to the radio). Therefore, detection of mm flux at a few hours post event should  be a good method of indicating a high redshift explosion.  ALMA, with its high sensitivity ($\\sim$75 $\\mu$Jy in 4 min), will be a potential tool for selecting potential high-$z$ bursts that would be high priority for intense followup across the spectrum. This will hopefully greatly increase the rate at which high-$z$ events are identified.  Finally, our data does not rule out a late jet break at $t_j>45$ d, which, as discussed above, makes the total explosion energy uncertain. Extremely sensitive VLA observations would be required to distinguish between the isotropic versus jet model. However, in 2010 with an order of magnitude enhanced sensitivity the EVLA will be the perfect instrument for such studies. For a  2\\,hr integration in 8 GHz band, the EVLA can reach sensitivity up to 2.3 $\\mu$Jy which will be able to detect the \\event\\ for 2 years or 6 months if the burst is isotropic or jet-like, respectively. EVLA will thus be able detect fainter events and follow events like GRB\\,050904 and \\event\\ for a longer duration, therefore obtaining  better density measurements, better estimates of outflow geometry and the total kinetic  energy."
    },
    "0910/0910.3771_arXiv.txt": {
        "abstract": "Spectral energy distributions (SEDs) of the central few tens of parsec region  of  some of the nearest, most well studied,  active galactic nuclei (AGN) are presented. These genuine AGN-core SEDs, mostly from Seyfert galaxies,  are characterised by two main features: an IR bump  with the maximum in the 2 -10 $\\mu$m range, and an increasing  X-ray spectrum   with frequency in    the  1 to $\\sim$ 200 keV region. These  dominant features are common  to  Seyfert type 1 and 2 objects  alike. In detail, type 1 AGNs    are clearly distinguished from type 2s by their high spatial resolution SEDs: type 2  AGN exhibit a sharp drop shortward of 2 $\\mu$m, with  the optical to UV region being fully absorbed;  type  1s  show instead a gentle 2 $\\mu$m drop   ensued by a secondary, partially-absorbed optical to UV emission bump. On the assumption that the bulk of optical to UV photons  generated in these AGN  are   reprocessed  by dust and re-emitted in the IR in an isotropic manner, the IR bump luminosity  represents  $\\gtrsim 70\\%$   of the total energy output   in these objects, the second energetically important contribution  are the high energies   above 20 keV. \\\\   Galaxies selected by their warm infrared colours, i.e. presenting a relatively-flat flux distribution in the 12 to 60 $\\mu$m range  have often being classified as   active galactic nuclei (AGN). The results from these high spatial resolution SEDs question this criterion as a general rule.  It is found that  the intrinsic  shape  of the infrared  spectral energy distribution of an AGN and inferred bolometric  luminosity  largely   depart from those derived  from   large aperture data.  AGN  luminosities can be overestimated by up to two orders of magnitude if relying on  IR satellite data. We find these differences  to be critical    for AGN luminosities below or about $10^{44}$ erg~ s$^{-1}$. Above this limit,     AGNs tend to dominate the light of their host galaxy regardless of the  integration aperture size used. Although the number of objects presented in this work is small,  we tentatively mark   this luminosity as a threshold to identify galaxy-light-  vs AGN- dominated objects. \\\\ ",
        "introduction": "}  The study of the spectral energy distributions (SED) over the widest possible spectral range is an  optimal way to characterise the properties of galaxies in general. Covering the widest spectral range is the key to  differentiate physical phenomena which  dominate at specific spectral ranges: e.g.  dust emission in the IR, stellar emission in the optical to UV, non-thermal processes in the X-rays and radio,  and to interrelate them as most of these  phenomena  involve radiation  reprocessing from a spectral range into another. The availability of the  SED of a galaxy   allows us to determine basic parameters such as its bolometric luminosity (e.g. Elvis et al. 1994;  Sanders \\& Mirabel 1996; Vasudevan \\& Fabian 2007), and via    modelling of the SED, its star formation level, mass  and  age (e.g. Bruzual \\& Charlot 2003; Rowan-Robinson et al. 2005;  Dale et al. 2007).  The construction of bona-fide SEDs is not easy  as it involves  data acquisition  from different ranges of the electromagnetic spectrum using very different telescope infrastructure. That already introduces a further complication as  the achieved spatial resolution, and with it the aperture size used,  vary with  the spectral range. SEDs based on the integration of the overall galaxy light may be very different from those extracted from only a specific region, for example the nuclear region. In this specific case,   the aperture size matters a lot, as different light sources, such as circumnuclear star formation, the active nucleus, and the subjacent galaxy light, coexist on small spatial scales and may contribute to the total nuclear output with comparable energies (e.g. Genzel et al.  1998; Reunanen et al. 2009).  In the specific case  of SEDs of AGN, it is often assumed that the AGN light dominates the integrated light of the galaxy at  almost any spectral range and for almost any aperture. This assumption becomes mandatory  at certain spectral ranges, such as  the high energies, the extreme UV or the mid-to-far-IR, because of  the spatial resolution limitations  imposed by the available instrumentation,  which currently lies in the several arcsecs  to  arcminutes  at these wavelengths. In the mid- to far- IR in particular, the available data, mostly  from IR satellites, are limited to spatial resolutions of a few arcsecs at best. Thus, the associated SEDs include  the contribution of  the host galaxy, star forming regions, dust emission and the AGN, with the first two components being measured over different spatial scales in the galaxy depending on the object distance and the spatial resolution achieved at a given IR wavelength.  In  spectral ranges where high spatial resolution is readily available,  the importance, if not dominance, of circumnuclear star formation relative to that of the AGN has become  clear  in the UV to optical range (e.g. Munoz-Marin et al. 1997), or in the near-IR  (Genzel et al. 1998). In the radio regime, the comparison of low- and high- spatial-resolution maps shows  the  importance of the diffuse circumnuclear emission and the emission from  the jet components  with respect to that of the  core itself (e.g Roy et al., 1994; Elmouttie et al, 1998; Gallimore et al. 2004; Val, Shastri \\& Gabuzda, 2004). Even with  low resolution data a major concern shared by most works is the relevance of the host galaxy contribution to the nuclear integrated emission  from the UV to optical to IR.  To overcome these mixing effects introduced by poor  spatial resolution,  different   strategies or assumptions have been followed  by the community. In quasars, by their own nature,  the dominance of the   AGN light over the integrated galaxy light at almost any wavelength is assumed; conversely, in lower luminosity AGN,   the contribution of different components is assessed via modelling of the integrated light (Edelson \\& Malkan 1986;  Ward et al. 1987; Sanders et al. 1988; Elvis et al. 1994; Buchanan et al. 2006 among others).  In this paper,  we attempt to provide a best estimate of  the AGN light  contribution on very nearby AGN by using    very high  spatial resolution data  over a wide range of the electromagnetic spectrum. Accordingly,  SEDs  of the central few hundred parsec  region of some of the nearest and brightest AGN  are compiled.   The work is motivated by the current possibility to obtain subarcsec resolution data in the near-to-mid-IR of  bright   AGN, and thus at  comparable resolutions to those available with  radio interferometry and  the HST in the optical to UV wavelength range. This is possible thanks to the use of adaptive optics in the near-IR, the diffraction limit resolutions provided by  8 -10 m telescopes in the mid-IR as well as  interferometry in the mid-IR.  The selection of targets is driven by the requirements imposed by the use of adaptive optics in the near-IR, which limits the observations to the availability of having bright point-like targets with magnitudes   V$~< $15 mag in the field, and the current flux detection limits in  mid-IR ground-based observations. AGN in the near universe are sufficiently bright  to  satisfy those criteria. The  near- to mid- IR high resolution  data used in this work come mostly from the ESO Very Large Telescope (VLT), hence this study relies on Southern targets, all well known objects, mostly Seyfert galaxies: Centaurus A, NGC 1068, Circinus, NGC 1097, NGC 5506, NGC 7582, NGC 3783, NGC 1566 and NGC 7469. For comparison purposes, the SED of the quasar 3C 273 is also included.  The compiled SEDs make use of the highest spatial resolution data available with current instrumentation across the electromagnetic spectrum. The main sources of data include: VLA-A array and ATCA data in radio, VLT diffraction-limited images and VLTI interferometry in the mid-infrared (mid-IR),   VLT adaptive-optics images in the near-infrared,  and \\textit{HST} imaging and spectra in the optical-ultraviolet. Although  X-rays and $\\gamma$-rays do not provide such a fine resolution,  information  when available for these galaxies  are also included in the SEDs on the assumption that above 10 keV or so we are sampling the  AGN core region. Most of the data used comes from the Chandra and INTEGRAL telescopes.  The novelty in the analysis  is the spatial resolutions achieved  in the infrared (IR), with  typical  full-width at half-maximum (FWHM) $\\lesssim $ 0.2 arcsec in the 1--5 $\\mu$m, $< $0.5 arcsec in the 11 -- 20 $\\mu$m. The availability of IR images  at  these spatial resolutions allow us to pinpoint the true spatial  location of the  AGN -- which happens not to have an  optical counterpart in most of the type 2 galaxies studied --  and extract its luminosity within aperture diameters of  a few tens of   parsec. The new compiled SEDs are presented in sect. 3. Some major differences but also similarities between the SEDs of  type 1 and type 2 AGN  arise at these resolutions.  These are presented and discussed in sections 4 and 5. The    SEDs and the inferred nuclear luminosities  are further compared with those extracted in the mid-to-far IR from large aperture data from IR satellites, and the differences discussed in sect. 6.  Throughout this paper, $H_0 = 70$ km s$^{-1}$ Mpc$^{-1}$ is used. The central wavelength of the near-IR broad band filters used are: $I$-band (0.8 $\\mu$m), $J$-band (1.26 $\\mu$m), $H$-band (1.66 $\\mu$m), $K$-band (2.18 $\\mu$m), $L$-band (3.80 $\\mu$m) and $M$-band (4.78 $\\mu$m). \\\\ ",
        "conclusions": "Sub-arcsec resolution data spanning the UV, optical, IR and  radio   have been used to construct  spectral energy distributions of the central, several  tens of parsec,  region of  some of the  nearest and brightest active galactic nuclei. Most of these objects are Seyfert galaxies. \\\\  These high spatial resolution SEDs differ largely from those derived from large aperture data, in particular in the  IR: the shape of the SED is different and the true AGN luminosity can get overestimated by orders of magnitude if based on IR satellite data. These differences appear to be critical for AGN luminosities below $10^{44} erg~s^{-1}$ in which  case large aperture data  sample  in full the host galaxy light.  Above that limit we  find cases among these nearby Seyfert galaxies where the  AGN   behaves  as the  most powerful  quasars,   dominating   the host galaxy light  regardless of the integration aperture-size used. \\\\  The high spatial resolution SED of these nearby AGNs  are all characterised by two major features in their power distribution: an IR bump  with maximum in the 2 -10  $\\mu$m range, and an increasing trend in X-ray power  with frequency in the  1 to $\\sim$ 200 keV region, i.e. up to the hardest energy that was possible to sample. These  dominant features are common  to  Seyfert type 1 and 2 objects  alike. \\\\  The major difference between type 1  and 2 in these SEDs arises  shortward of 2 $\\mu$m. Type 2s are  characterised by a sharp fall-off shortward of this wavelength, with no optical counterpart to the IR nucleus being detected  beyond  1 to 0.8 $\\mu$m. Type 1s show   also a drop shortward of 2 $\\mu$m but this is   more gentle - the spectrum is flatter -  and  recovers at about 1 $\\mu$m to give rise to the characteristic   blue-bump feature seen in quasars. The flattening of the spectrum   shortward of 2 $\\mu$m is also an expected feature  of type 1 AGNs. Interpreting  the  IR bump  as AGN reprocessed emission by the nuclear dust,  in type 1s   the nearest to the centre hotter dust can be directly seen, hence the flattening of their spectrum, whereas in type 2s this hot dust   is still fully obscured. \\\\  Longward of 2 $\\mu$m, all the AGN types   show very similar SEDs, the bulk of the IR emission starts from this wavelength on and  the shape of the IR bump is very similar in all the AGNs. This is compatible with an equivalent black-body temperature for the bulk of the  dust in the 200 - 400 K range in average. Although the current shape of the IR bump is limited by the availability of high angular resolution data  beyond    20  $\\mu$m  for most objects, due to  the small   region sampled in these SEDs, of just a  few parsecs in some galaxies, a major contribution from colder dust that will modify the IR bump is not expected. \\\\  It can thus be concluded that at the scales of  a few tens of parsec from the central engine, the  bulk of the IR emission in either AGN type can be reconciled with pure dust emission. It follows that further    contributions from   a non-thermal synchrotron component and/or  a thermal free-free emission linked to cooling of ionised gas  are insufficient to overcome that of  dust at these physical scales. The detailed modelling of  NGC 1068's SED - this being one of the most complete we have compiled -  in which   these three contributions --   synchrotron, free-free and a dust torus   components -  are  taking into account illustrates that premise, that is, the dominance of dust emission in the IR, even at the parsec-scale resolution achieved for  this object in the mid-IR with  interferometry (Hoenig et al.  2008). Only the two more extreme objects in this analysis, Cen A, on the low luminosity rank, and  3C 273, on the highest, present a SED that is not   dominated by dust but by  a  synchrotron component. We tend to  believe that is due to a much reduced dust content  in these nuclei. \\\\   Over the nine orders of magnitude in frequency covered by these SEDs, the power stored in the IR bump is by far the most energetic  fraction of the total energy budget measured in these objects.  Evaluating this total budget as  the sum of the  IR  and hard X-ray -- above 20 keV --  luminosities,  the IR part accounts for more than 70\\% of the this total  in  seven out of the ten AGN studied. In the three exceptions, the IR fraction reduces to $<\\sim 30\\% $ (3C 273 and NGC 1566), $< \\sim 60\\%$ in Cen A. Even if accounting for variability in the X-rays, by a factor 2 to 3 in average, the IR emission    remains in all cases dominant  over, or as  important as, in the last three cases, the  X-ray emission. If comparing with the  observed blue bump luminosity   in the type 1 nuclei, this  represents less than  15\\% of the IR emission. Putting all together, the IR bump  energy from these high spatial resolutions SEDs  may represent the tightest measurement of the accretion luminosity in these Seyfert AGN. \\\\   The average high spatial resolution SED  of the type 2 and  of the  type 1 nuclei analysed in this work, and presented in Fig. 4, can be retrieved from http://www.iac.es/project/parsec/main/seyfert-SED-template. \\\\  This work was initiated and largely completed during the stay of   K. Tristram, N. Neumayer and A. Prieto at the Max-Planck Institut fuer Astronomie in  Heidelberg."
    },
    "0910/0910.1774_arXiv.txt": {
        "abstract": "We present \\textit{B} and \\textit{R} band spectroastrometry of a sample of 45 Herbig Ae/Be stars in order to study their binary properties. All but one of the targets known to be binary systems with a separation of $\\rm \\sim0.1-2.0$~arcsec are detected by a distinctive spectroastrometric signature. Some objects in the sample exhibit spectroastrometric features that do not appear attributable to a binary system. We find that these may be due to light reflected from dusty halos or material entrained in winds. We present 8 new binary detections and 4 detections of an unknown component in previously discovered binary systems. The data confirm previous reports that Herbig Ae/Be stars have a high binary fraction, $\\rm{74\\pm6}$~per cent in the sample presented here. We use a spectroastrometric deconvolution technique to separate the spatially unresolved binary spectra into the individual constituent spectra.  The separated spectra allow us to ascertain the spectral type of the individual binary components, which in turn allows the mass ratio of these systems to be determined. In addition, we appraise the method used and the effects of contaminant sources of flux. We find that the distribution of system mass ratios is inconsistent with random pairing from the Initial Mass Function, and that this appears robust despite a detection bias. Instead, the mass ratio distribution is broadly consistent with the scenario of binary formation via disk fragmentation. ",
        "introduction": "Our understanding of the formation and early evolution of massive stars ($\\rm M_* \\ga 8M_{\\odot}$) is much less complete than in the case of low mass stars. The scenario of low mass star formation has been relatively well studied, and a broadly consistent observational and theoretical picture has now emerged. The various phases of low mass star formation include: cloud collapse, proto-stellar creation and a subsequent contraction of Pre Main Sequence (PMS) objects towards the Zero Age Main Sequence (ZAMS). This later stage, the T Tauri phase, is easy to observe and therefore relatively well understood \\citep{Bouvier2007}. In the case of more massive stars the situation is much less clear. Such stars do not experience an optically visible PMS phase, evolve on a much more rapid timescale, and are considerably more luminous than low mass stars. Early studies on the effects of radiation pressure and the considerable ionising output of massive young stars prompted speculation that massive star formation might proceed in a different manner to that of low mass stars \\citep{Larson1971, Kahn1974}. For example, it has been suggested that the most massive stars form via stellar mergers or competitive accretion \\citep{JBally2005}.  \\smallskip  However, recent work, on both the observational and theoretical front, suggests that massive star formation may not be dissimilar to low mass star formation. As an example of observational results, \\citet{Pateletal2005} report the detection of a massive disk around a 15$\\rm M_{\\odot}$ protostar, indicating that massive stars may form via monolithic accretion. On the theoretical front, recent work indicates accretion onto a massive protostar is not impeded by radiation pressure \\citep{YorkeandSonnhalter2002,Turner2007,Krumholz2009}. However, while significant progress has been made, there remain many unaddressed questions related to the formation and evolution of massive stars \\citep{ZinneckerandYorke2007}. As observations of massive young stars are challenging, the full extent of the differences and similarities between low and high mass star formation are still unknown.  \\smallskip   Between the two extremes of mass lie the Herbig Ae/Be (HAe/Be) stars \\citep{Herbig1960}. These stars represent the most massive of objects to experience an optically visible PMS evolutionary phase. Therefore, HAe/Be stars offer an opportunity to study the early evolution of stars more massive than the sun. Spectropolarimetry indicates that Herbig Ae stars may undergo a PMS phase similar to that of the T Tauri stars, while Herbig Be stars may evolve via disk accretion, rather than magnetospheric accretion \\citep{JSVink2002,JSVink2005a,JCMottram2007}. Therefore, it appears that a transition in formation mechanisms occurs across the HAe/Be mass boundary \\citep{JCMottram2007}. However, the critical mass has not yet been established.  \\smallskip  To examine the similarities and differences between low mass T Tauri stars, HAe/Be stars and the optically invisible Massive Young Stellar Objects (MYSOs), study of the circumstellar environment at small angular scales is required. This is not trivial, requiring observations with high angular resolution \\citep{Mannings1997,Fuente2006,Grady2007,Kraus2008}. Despite the progress in the field, a full understanding of HAe/Be stars is hampered by the small sample sizes involved. By way of contrast, \\citet{DB2006} utilised spectroastrometry to study a large sample of HAe/Be stars with milli-arcsecond (mas) precision. Despite this resolution \\citet{DB2006} did not detect any accretion disks around HAe/Be stars. However, they did find that the majority, $\\rm{68\\pm11}$~per cent, of HAe/Be stars reside in relatively wide (probably a few-hundred~au, see Section \\ref{sep}) binary systems.  \\smallskip  The binary fraction reported by \\citet{DB2006} is greater than that of T Tauri stars at similarly wide separations, which in turn is greater than that of Main Sequence G-dwarfs at the same separations \\citep{DuquennoyandMayor1991,Reipurth1993,Ghez1993}. Indeed, this high binary fraction is approaching that of more massive stars \\citep{Preibisch1999}. However, little is known about the properties of such binary systems. The properties of the binary components and configurations of such systems are of interest as they can constrain the binary formation mechanism. The seminal study to date is that by \\citet{BC2001}, who used Adaptive Optics assisted observations to construct Spectral Energy Distributions (SEDs) for each component in a number of HAe/Be binary systems. The drawback of SED fitting is that PMS stars, as young stars, are inevitably associated with dusty, obscured environments. Therefore, the brightness ratio of a binary determined by SED fitting can occasionally be ambiguous. However, very few HAe/Be binary systems have been studied with spatially resolved spectroscopy, and thus far such studies have been conducted with seeing limited resolution \\citep{ACarmona2007,Hubrig2007}.  \\smallskip   The position angles of HAe/Be binary systems seem to be preferentially aligned with the spectropolarimetrically detected circumprimary disks \\citep{DB2006}. This already places constraints on the formation modes of these stars, in that it seems the systems formed via fragmentation of a molecular core or disk. This had already been suggested for lower mass binaries \\citep{SWolf2001,Kroupa2001}, but little is known about the formation mechanisms of more massive stars. This paper describes a spectroastrometric follow-up of the work of \\citet{DB2006} with dedicated observations to study both components of binary systems. The objective is to determine the properties of these binary systems and thus place stronger, more quantitative, constraints on the formation of stars of intermediate mass. We do this by determining the mass ratio of these binary systems. This is done using a spectroastrometric technique to disentangle the constituent spectra of unresolved binary systems, allowing the spectral type, and hence mass, of each component to be determined. Spectroastrometry itself is a relatively simple technique that extracts the spatial information present in conventional longslit spectra. Crucially, spectroastrometry can probe changes in flux distributions with a typical precision of a mas or less \\citep{JBailey1998a}, which is required to study unresolved binary systems. Typically the minimum separation probed is of the order 100~mas, as the signature of a binary system is dependant upon the system brightness ratio and separation.  \\smallskip  This paper is structured as follows: in Section \\ref{obsanddatred} we present our sample selection, observation method and data reduction procedures. In Section \\ref{results_spec_ast} we discuss the spectroastrometric signatures observed. In Section \\ref{spec_split_app} we present the method of splitting unresolved binary spectra and in Section \\ref{results_spec} we review the results of separating binary spectra into their constituent spectra. In Section \\ref{discussion} we discuss our results. Finally, we conclude this paper in Section \\ref{conclusions} by summarising the salient points raised.  \\section[]{Observations and data reduction} \\label{obsanddatred} \\subsection{Observations}  The data presented consist of long-slit spectra in the \\textit{B} band ($\\rm 4200-5000 \\rm{\\AA}$) and/or the \\textit{R} band ($\\rm 6200-7000 \\rm{\\AA}$) of 45 HAe/Be stars, and 2 emission line objects which are possible HAe/Be stars. The objects were chosen from the catalogs of \\citet{Theetal1994},\\,\\citet{Vieira2003}\\,\\&\\,\\citet{JHernandez2004}, and were selected to be reasonably bright (\\textit{V} $\\rm \\leq $ 12-13). Some objects previously observed by \\citet{DB2006} were observed to provide a consistency check on the spectroastrometric signatures. Given the small population of HAe/Be stars, the objects observed constitute a representative sample of HAe/Be stars, albeit brightness limited.  \\smallskip  The data were obtained using the 4.2m William Herschel Telescope (WHT) and the 2.5m Isaac Newton Telescope (INT). At the WHT, data were obtained on the 6th \\& 7th of October 2006, using the Intermediate Dispersion Spectrograph and Imaging System (ISIS) spectrograph. Spectra of 20 objects were taken simultaneously in the \\textit{B} and \\textit{R} bands using the dichroic slide of ISIS. In most cases a slit $\\rm{5}$~arcsec wide was used to ensure all the light from a given binary system entered the slit, even in poor seeing. This allows us to study the individual binary components, unlike \\citet{DB2006}, who used a slit of 1~arcsec. The R1200B and R1200R gratings were used and the resulting spectral resolving power was found to be $\\rm \\sim 3500$, corresponding to $\\rm 85\\, km\\,s^{-1}$. The angular pixel size was $\\rm{0.20}$ and $\\rm{0.22}$~arcsec in the \\textit{B} band and \\textit{R} band respectively, which means that the spatial profile of the longslit spectra was well sampled (average FWHM 1.9~arcsec). At the INT data were obtained using the 235mm camera and the Intermediate Dispersion Spectrograph (IDS). Observing was conducted from the 27th of December 2008 to the 3rd of January 2009. The spectra of 32 objects were obtained, despite adverse weather conditions preventing observing for the better part of three nights. As at the WHT the slit width was generally $\\rm5$~arcsec. The R1200R and R1200B gratings were used and the resulting spectral resolution was found to be $\\rm \\sim 3800 $, or $\\rm80 \\,km\\,s^{-1}$. The angular size of the pixels was $\\rm{0.4}$~arcsec, which fully sampled the average spatial profile of the spectra (1.8 arcsec).  \\smallskip  Multiple spectra were taken at four position angles (PA) on the sky. The PAs selected always comprised of two perpendicular sets of two anti-parallel angles, e.g. $\\rm 0\\degr$, $90^{\\circ}$, $180^{\\circ}$ and $270^{\\circ}$. Dispersion calibration arcs were made using CuNe and CuAr lamps. Table \\ref{logofobs} presents a summary of the observations.   \\begin{center} \\begin{table*} \\begin{minipage}{\\textwidth} \\caption{\\label{logofobs}\\small{Log of the observations, column 1 lists the objects observed, column 2 denotes the spectral type of the objects, column 3 lists the \\textit{V} band magnitudes of the sample, and column 4 designates which telescope the object in question was observed with. Columns 5 and 6 list the average seeing conditions, columns 7 and 8 list the total exposure times and column 9 denotes the slit width used. Column 10 lists the total Signal to Noise Ratios, and finally, column 11 presents the date(s) each object was observed. Information on the objects is taken from SIMBAD (simbad.u-strasbg.fr) unless otherwise stated.}} \\begin{center}  \\begin{tabular}{p{1.725cm} p{1.25cm} p{0.7cm} p{0.95cm} p{0.90cm}  p{0.90cm} p{0.550cm} p{0.550cm} p{1.1cm} p {1.35cm} p{3.15cm}} \\hline Object & Spec type & \\hspace*{2mm}\\textit{V}  &Telescope & $\\rm \\overline{FWHM}^a$ & $\\rm \\overline{FWHM}^b$& $\\rm{t_{blue}}$ &$\\rm{t_{red}}$ & Slit & SNR & Date\\\\ &  &    &      &(\\arcsec)& (\\arcsec)   & (s)  & (s) & (\\arcsec) \\\\ \\hline \\hline VX Cas & $\\rm{A0e}$  & 11.3 &  WHT & 1.3 &1.2  & 4800 & 4800  & 5.0&$\\rm{600_B}$,$\\rm{570_R}$ &07/10/2006 \\\\ VX Cas  & $\\rm{A0e}$  & 11.3 &  INT & 1.1 & 1.7 & 2800 &  3600 & $\\rm 2.5_{\\rm B},5.0_{\\rm R}$ & $\\rm{370_B}$,$\\rm{370_R}$&28/12/2008,31/12/2008\\\\ V594 Cas    & Be & 10.6 & INT &1.3 & -- & 3200 &-- &5.0 &610 &01/01/2009\\\\ V1185 Tau  & A1 & 10.7 & INT&  1.7 & --& 3200 &-- & 5.0 & 430&03/01/2009 \\\\ IP Per  &  $\\rm{A3}$ & 10.3 & INT & 1.2 & 1.4 &2000 & 2400 &$\\rm 2.5_{\\rm B},5.0_{\\rm R}$ & $\\rm{110_B}$,$\\rm{320_R}$&28/12/2008,31/12/2008\\\\ AB Aur & A0Vpe & 7.1  & WHT & 1.9 & 1.9 & 330  & 320  & 5.0  &$\\rm{110_B}$,$\\rm{650_R}$ &06/10/2006\\\\ MWC 480 &   $\\rm{A3pshe}$& 7.7 &  WHT & 2.0 & 2.1 & 960  & 640 & 5.0 & $\\rm{1100_B}$,$\\rm{800_R}$&06/10/2006\\\\ UX Ori &  $\\rm{A3e}$  & 9.6 &  WHT & 2.4 &2.4 & 3600 & 3600 & 5.0 & $\\rm{1200_B}$,$\\rm{940_R}$&06/10/2006\\\\ V1012 Ori   & $\\rm{Be^c}$ &12.1  &INT &2.1& -- &  4800 & --& 5.0 &150 &02/01/2009\\\\ V1366 Ori    & A0e& 9.8& INT& 1.3 &--  & 2400 & --& 3.0 &570 &31/12/2008 \\\\ V346 Ori  &   $\\rm{A5III}$&10.1 & INT & 1.5 & -- & 3600 &-- & 5.0 & 200&01/01/2009\\\\ HD 35929  & A5 & 8.1 &   WHT & 1.9  & 1.7 & 2060 & 1470 & 5.0 & $\\rm{40_B}$,$\\rm{900_R}$&07/10/2006\\\\ V380 Ori &  A0& 10.7 & INT  & 1.5 &1.6 & 3600 & 2940 & $\\rm 3.0_{\\rm B},5.0_{\\rm R}$& $\\rm{100_B}$,$\\rm{200_R}$&28/12/2008,31/12/2008\\\\ MWC 758 &   A3e & 8.3 &  WHT & 1.4 & 1.3&1080 & 960& 5.0  &$\\rm{50_B}$,$\\rm{660_R}$ &07/10/2006\\\\ HK Ori  & A4pev &11.9 & INT & 2.4 & -- & 4800 & --& 5.0 & 200&02/01/2009\\\\ HD 244604  & A3 & 9.4 &  WHT & 1.7  & 1.6 & 3180 & 3660 & 5.0 &$\\rm{100_B}$,$\\rm{720_R}$ & 07/10/2006\\\\ V1271 Ori  &  A5 & 10.0  & INT &1.6&-- & 2460 & --& 5.0 & 410& 01/01/2009\\\\ T Ori      & A3 & 9.5 & INT &1.9 & -- & 3200& -- & 5.0 & 300 &03/01/2009 \\\\ V586 Ori & $\\rm{A2V}$ &9.8 & INT & 3.3 & -- & 2940  & --& 5.0 & 650&02/01/2009\\\\ HD 37357 & A0e &8.8 & INT &  1.4 & 1.4 &2060 & 1470& $\\rm 3.0_{\\rm B},5.0_{\\rm R}$ & $\\rm{200_B}$,$\\rm{350_R}$ & 28/12/2008,31/12/2008\\\\ V1788 Ori  &  B9Ve & 9.9 & INT& 1.7 &-- & 1350  &-- & 5.0 & 450 &01/01/2009\\\\ HD 245906   & B9IV & 10.7 & INT & 1.8 & --  &2800  & --& 5.0  &100 &03/01/2009\\\\ RR Tau       & A2II-IIIe & 10.9& INT & 1.7 & -- &2800 &-- & 5.0 & 250&03/01/2009 \\\\ V350 Ori     & A0e & 10.4(\\textit{B})& INT &  1.9 & -- &4800 &-- & 5.0 & 130& 03/01/2009 \\\\ MWC 120 &  A0 & 7.9 &  WHT & 2.1  & 1.9 & 480 & 480 &5.0 &$\\rm{1500_B}$,$\\rm{690_R}$ & 06/10/2006\\\\ MWC 120  &  A0 & 7.9 &  INT & 1.4& 1.6  & 2460 & 1250 & $\\rm 3.0_{\\rm B},5.0_{\\rm R}$ & $\\rm{1200_B}$,$\\rm{560_R}$&28/12/2008,31/12/2008\\\\ MWC 790     & Be & 12.0 & INT & 3.1 & --  & 4050  &-- & 5.0 &200 & 02/01/2009 \\\\ MWC 137 & Be & 11.2& INT& -- &  2.1 &-- & 4560& 5.0 & 200 & 28/12/2008 \\\\ HD 45677 &  $\\rm{Bpshe}$ & 8.0  & WHT &  2.0  & 1.9 &360  &240  &5.0 & $\\rm{900_B}$,$\\rm{550_R}$& 06/10/2006\\\\ LkH$\\rm \\alpha$ 215 &   B7.5e  & 10.6 & INT & 1.5 & 2.3 &3600 & 3600& $\\rm 5.0_{\\rm B},4.0_{\\rm R}$ & $\\rm{300_B}$,$\\rm{360_R}$ & 27/12/2008,31/12/2008\\\\ MWC 147 &   $\\rm{B6pe}$ & 8.8 & WHT & 1.8 & 1.5 & 3000 & 1700 & 5.0 & $\\rm{1200_B}$,$\\rm{500_R}$& 07/10/2006\\\\ MWC 147   &  $\\rm{B6pe}$ & 8.8 & INT & 1.3 & --  & 2400 & --& 5.0 & 620& 01/01/2009\\\\ R Mon &  B0 & 10.4& INT & 4.2 & 2.5  & 3600 & 3600 & 5.0& $\\rm{100_B}$,$\\rm{240_R}$& 28/12/2008,02/01/2009\\\\ V590 Mon  &  B8pe & 12.9  & INT & 1.5 & -- & 2670 & --& 5.0 & 200 & 01/01/2009\\\\ V742 Mon  & B2Ve & 6.9 & INT & 1.4  & 2.8& 1740 & 2535 & 5.0 & $\\rm{400_B}$,$\\rm{800_R}$ & 30/12/2008,31/12/2008\\\\ OY Gem       & Bp[e] & 11.1 & INT &1.8 & -- & 2880 & --& 5. 0 & 100 & 03/01/2009\\\\ GU CMa&  $\\rm{B2Vne}$ & 6.6  & WHT & 2.5  & 2.4 & 360 & 360 & 5.0 & $\\rm{1500_B}$,$\\rm{900_R}$&06/10/2006 \\\\ GU CMa  &  $\\rm{B2Vne}$ & 6.6  & INT &1.8 & -- & 720 &--& 5.0 & 1400 &03/01/2009 \\\\ MWC 166 &  B0IVe & 7.0  & WHT &  2.5 & 2.3& 210& 120  &5.0 & $\\rm{1200_B}$,$\\rm{800_R}$& 06/10/2006\\\\ HD 76868 & B5  & 8.0 & INT &1.5 & 2.4 & 4830 & 2100& 5.0 & $\\rm{100_B}$,$\\rm{100_R}$ & 30/12/2008,01/01/2009\\\\ HD 81357   & B8 & 8.4 & INT & 1.7 & --  & 4800 & -- & 5.0& 100 &03/01/2009\\\\ MWC 297 & $\\rm{Be}$  & 12.3 & WHT &  1.4 &  1.2 &4100 & 3120 &5.0 & $\\rm{100_B}$,$\\rm{500_R}$&07/10/2006\\\\ HD 179218 &  B9e  & 7.2  & WHT & 2.6 & 2.4 &  2100& 1200 &1.0/1.5 & $\\rm{2300_B}$,$\\rm{940_R}$&06/10/2006\\\\ HD 190073 & $\\rm{A2IVpe}$ & 7.8  & WHT & 1.5 & 1.2 &540  & 360&5.0 &$\\rm{600_B}$,$\\rm{800_R}$ &07/10/2006\\\\ BD +40 4124&B2 & 10.7  & WHT & 2.1 & 1.9 &  600 &  660 & 4.0 & $\\rm{500_B}$,$\\rm{370_R}$ &07/10/2006\\\\ MWC 361  &  $\\rm{B2Ve}$ & 7.4  & WHT & 1.7  & 1.7 & 1350 & 960 &2.5/4.0 & $\\rm{1400_B}$,$\\rm{1400_R}$ &06/10/2006\\\\ SV Cep  &  $\\rm{Ae}$& 10.1(\\textit{B}) & INT & 1.6& --& 3000& -- & 5.0 & 600 &02/01/2009\\\\ MWC 655     &  B1IVnep & 9.2 & INT & 1.7 & --  & 2400 & -- & 5.0 & 400&03/01/2009\\\\ Il Cep  &  $\\rm{B2IV/Ve}$  & 9.3  & WHT & 1.4 & 1.2 &3500 & 3000&5.0& $\\rm{800_B}$,$\\rm{500_R}$&07/10/2006\\\\ BHJ 71&B4e & 10.9 &  WHT &  1.8  & 1.8 & 1200 & 1080 &4.0 & $\\rm{500_B}$,$\\rm{340_R}$ &06/10/2006\\\\ BHJ 71  & B4e & 10.9 &  INT &  1.7 & -- &4200 & --& 5.0 & 500 &01/01/2009\\\\ MWC 1080 &  $\\rm{B0}$ & 11.6 &  WHT & 2.0 & 2.0 &  3300 & 4170 & 5.0 &$\\rm{200_B}$,$\\rm{500_R}$& 06/10/2006\\\\ \\hline  \\end{tabular}  \\end{center} \\renewcommand\\footnoterule{} \\footnotetext{{$\\rm^a$ Average seeing in the blue spectral region, approximated by the average of the individual median FWHM, where necessary averaged\\\\ \\hspace*{3.4mm} over multiple slit widths.}} \\footnotetext{{$\\rm^b$ Average seeing in the red spectral region, approximated by the average of the individual median FWHM, where necessary averaged \\\\ \\hspace*{3.4mm}over multiple slit widths.}} \\footnotetext{{$\\rm^c$ \\citet{Theetal1994}}} \\end{minipage} \\end{table*} \\end{center}   \\subsection{Data reduction} \\label{data_red}    Data reduction was conducted using the Image Reduction and Analysis Facility (IRAF)\\footnote{IRAF is distributed by the National Optical Astronomy Observatories, which are operated by the Association of Universities for Research in Astronomy, Inc., under cooperative agreement with the National Science Foundation \\citep{IRAF}.} and routines written in Interactive Data Language (IDL). Initial data reduction consisted of bias subtraction and flat field division. The total intensity spectra were then extracted from the corrected data in a standard fashion. Wavelength calibration was conducted using the arc spectra, and the wavelength calibration solution had a precision of the order $\\rm{<0.1\\AA}$.  \\smallskip  Spectroastrometry was performed by fitting Gaussian functions to the spatial profile of the long-slit spectra at each dispersion pixel. This resulted in a positional spectrum, the centroid of the Gaussian as a function of wavelength, and a Full Width at Half Maximum (FWHM) spectrum, the FWHM as a function of wavelength. Spot checks were used to ensure that a Gaussian was an accurate representation of the data. The continuum position exhibited a general trend across the CCD chip: of the order of 10 pixels in the case of the ISIS data and 2 pixels in the IDS data. This was removed by fitting a low order polynomial ($\\rm{4^{th}}$ or $\\rm{5^{th}}$ order) to the continuum regions of the spectrum.  \\smallskip  All intensity, positional and FWHM spectra at a given PA were combined to make an average spectrum for each PA. A correction for slight changes in the dispersion across PAs was determined by cross-correlating average intensity spectra obtained at different PAs. The correction was then applied to the average intensity, positional and FWHM spectra.  The average positional spectra for anti-parallel PAs were then combined to form the average, perpendicular, position spectra, for example: ($\\rm 0^{\\circ} - 180^{\\circ}$)/2 and ($\\rm 90^{\\circ} - 270^{\\circ}$)/2. This procedure eliminates instrumental artifacts as real signatures rotate by $\\rm 180^{\\circ}$ when viewed at the anti-parallel PA, while artifacts remain at a constant orientation. In addition, all positional spectra were visually inspected for artifacts not fully removed by this procedure. As with the positional spectra, the FWHM spectra at anti-parallel PAs were also combined to make two averaged, perpendicular spectra. While FWHM features do not rotate across different PAs, the features observed at anti-parallel PAs were used to exclude artifacts via a visual comparison. All conditions being constant, a real FWHM signature should not change from one PA to the opposite angle at $\\rm +180^{\\circ}$. ",
        "conclusions": "\\label{conclusions}  In this paper we present spectroastrometric observations of a relatively large sample of HAe/Be stars. Here we present the salient findings of this work:  \\begin{itemize}  \\item{We find a high binary fraction, $\\rm{74\\pm6}~per cent$, consistent with previous studies.}  \\item{Using spectroastrometry to separate the unresolved binary spectra we determine spectral types for the components of 9 systems.}  \\item{The mass ratios of these systems, determined from the constituent spectral types, are inconsistent with a secondary mass randomly selected from the IMF.}  \\item{Although our sample is small this result constrains the mode of binary formation in that the mass ratios and separations of the binary systems observed suggest that the secondary forms via disk fragmentation.}  \\item{The properties of the binary systems observed indicate that these systems have not spent a significant amount of time in dense, clustered environments. Therefore, these systems demonstrate that isolated star formation can produce stars as massive as $\\rm{\\sim10-15M_{\\odot}}$.}  \\end{itemize}"
    },
    "0910/0910.1632_arXiv.txt": {
        "abstract": "\\begin{center} \\end{center}  We present a type of dark energy models where  the particles of dark energy $\\phi$  are dynamically produced via a quantum transition at very low energies. The scale where the transition takes places depends on the strength $g$ of the interaction between $\\phi$ and a relativistic field $\\vp$. We show that a $g\\simeq 10^{-12}$ gives a generation scale $E_{gen} \\simeq 1\\,eV$ with a cross section $\\sigma\\simeq 1 \\,pb$ close to the WIMPs cross section $\\sigma_{w}\\simeq pb$ at decoupling.  The number density $n_\\phi$  of the $\\phi$ particles is a source term in the  equation of  motion of $\\phi$ that generates the scalar potential $v(\\phi)$ responsible for the late time acceleration of our universe. Since the appearance of $\\phi$ may be at very low scales, close to present time, the cosmological coincidence problem can be explained simply due to the size of the coupling constant. In this context it is natural to unify dark energy with inflation in terms of  a single scalar field $\\phi$. We use the same potential $v(\\phi)$ for inflation and dark energy. However, after inflation $\\phi$ decays completely and reheats the universe at a scale $E_{RH} \\propto h^2 \\, m_{Pl}$, where $h$ is the coupling between the SM particles and $\\vp$. The field $\\phi$ disappears from the spectrum during  most of the time, from reheating until its  re-generation at late times, and therefore it does not interfere with the standard decelerating radiation/matter cosmological model allowing for a successful unification scheme. We show that the same interaction term that gives rise to the inflaton decay accounts for the late time re-generation of the $\\phi$ field giving rise to dark energy. We present a simple model where the strength of the $g$ and  $h$ couplings are set by the inflation scale $E_I$ with $g=h^2 \\propto E_I/m_{Pl}$ giving a reheating scale $E_{RH} \\propto E_I $ and $\\phi$-generation scale $E_{gen} \\propto E_I^2/m_{pl} \\ll E_{RH}$. With this identification we reduce the number of parameters and the appearance of dark energy is then given  in terms of the inflation scale $E_I$. ",
        "introduction": "The nature and dynamics of Dark Energy ''DE'', which gives the accelerating expansion  of the universe at present time, is now days one of the most interesting and stimulating fields of physics. It was discovered more than ten years ago \\cite{first aceleration} and it has been confirmed by further cosmological observations, being now one of the most robust conjecture  in modern physics. In fact the acceleration of the present universe has many experimental proofs such as  the CMB temperature and fluctuations \\cite{komatsu wmap}, in the matter power spectrum measured by galaxy surveys \\cite{sdss,Tegmark SDSS} and in type Ia supoernovae \\cite{Riess IA supernova,Kowalski supernova,Miknaitis supernova}.  The most popular model is the so called $\\Lambda$CDM model, in which a cosmological constant and some amount of cold dark matter are included ''by hand''. Despite  its extraordinary consistence with observations, $\\Lambda$CDM  is an effective model that leaves many unsolved theoretical question. In fact the existence of the cosmological constant and its order of magnitude have   no theoretical justification in  $\\Lambda$CDM. The cosmological coincidence problem, that is why the universe is starting to accelerate right now, is also unsolved. Introducing a cosmological constant at the initial stages of the standard cosmological model is specially troublesome since one has to fine tune its value to one parte in $10^{120}$. This problem can be ameliorated if we understand when and how dark energy appears in the universe and this is the main motivation of our present work \\cite{axelfab}. However, our approach will also help us to   unify dark energy with inflation.  An attractive dark energy alternative to the $\\Lambda$CDM model consists in introducing a ''quintessence'' scalar field $\\phi$ that generates the accelerating expansion \\ci{Q}\\ci{axQ} of the universe due to is dynamics. The dynamics is fixed by its potential $v(\\phi)$ and it is possible to choose potentials that lead to a late time acceleration of the universe \\ci{axQ}. This scalar field can be a fundamental  or composite particle as for example bound states \\ci{axDG}. In the second case, the bound quintessence fields are scalar fields composed of fundamental fermions, such as meson fields, and  can be generated at low energies as a consequence of a low phase transition scale due to a strong gauge coupling constant \\ci{axDG}. This allows to understand why DE appears at such late times. On the other hand, in the former case the appearance of fundamental scalar field is right at the beginning of the reheated universe and the acceleration of the universe takes place at a much later time due to the classical evolution of the quintessence field $\\phi$. The huge difference in scales between the reheating and dark energy scales requires a fine tuning in the choice of the potential.  Here, we present an interesting alternative, namely that the emergence of the fundamental quintessence particles $\\phi$ is originated from a quantum transition taking place at low energies, e.g. as low as $eV$  \\cite{axelfab}. The scale where this transition takes places depends on the strength of the interaction between $\\phi$ and a relativistic field $\\vp$ and it is dynamically determined by the ratio  $\\Gamma/H$ where $\\Gamma$ is the transition rate   and   $H$ the Hubble constant. A value of the coupling $g\\simeq 10^{-12}$ gives a generation scale $E_{gen}\\simeq 1\\,eV$ with a cross section $\\sigma = g^2/32 \\pi E_{gen} \\simeq 1 \\, pb$ close to the WIMPs cross section $\\sigma_{w}\\simeq O(pb)$ at decoupling\\ci{wimp}. The subsequent acceleration of the universe is due to the classical evolution of $\\phi$ due to the scalar potential $v(\\phi)$. Our quantum generation scheme does not aim to derive the potential $v(\\phi)$ but to understand why dark energy dominates at such a late time. Clearly, by closing the gap between the energy today $E_o$, where the subscript $o$ always refers to present time quantities,  and that of $\\phi$ production $E_{gen}$, we do not require a fine tuning of the parameters in $v(\\phi)$. Since the appearance of $\\phi$ may be at such low scales this offers a new interpretation and solution to the cosmological coincidence problem in terms of the size of the coupling constant $g$.  Furthermore, this late time production  of the $\\phi$ particles allows to implement in a natural way a dark energy-inflaton unification scheme. In this scenario, after inflation the field $\\phi$ decays completely and reheats the universe with standard model particles. The universe expands then in a decelerating  way dominated first by radiation and later by matter. At low energies the same interaction term that gives rise to the inflaton decay accounts for the quantum re-generation of the $\\phi$ field giving rise to dark energy. In general, it is not complicated to choose a scalar potential such that   the universe accelerates in two different regions, at early inflation and dark energy epochs, as in quintessential models \\ci{scalarunification}. However, the universe requires to be most of the time dominated first by standard model relativistic particles  and later by matter. The reheating of the universe and the long period of decelerating phase are  usually not taken into account in inflation-dark energy unification  models and these features are essential in the standard cosmological Big Bang model. In our case, the inflaton-dark energy field is completely absent during most of the time (from reheating until re-generation) and therefore it does not interfere with the standard cosmological model. We will exemplify our inflation-dark energy unified scheme with a simple model. The scalar potential $v(\\phi)$ will have only two parameters fixed by the conditions to give the correct density perturbations $\\delta \\rho/\\rho$ and the present time dark energy scale. The two couplings $h,g$, which give the strength of reheating process with SM particles and the $\\phi$ re-generation process at low energies, respectively, are free parameters but we may take them as $g=h^2 \\propto E_I/m_{pl}$, where $E_I$ is the scale of inflation and it is one of the parameters of $v(\\phi)$. Therefore, starting with four free parameters we can reduce the number to only two and these two are fixed by observations. This gives a reheating scale $E_{RH} \\propto E_I$ and $\\phi$-generation scale $E_{gen}\\propto E_I^2/m_{pl} \\ll E_{RH}$.  The paper is organized as follows: in section \\ref{motivations} we give an overview of the late time quantum generation of the quintessence field $\\phi$ and its possible unification with the inflaton field. In section \\ref{regeneration} we present the dark energy quantum generation in detail. In section \\ref{regeneration} we show how to unify inflation and dark energy in terms of a single scalar field in the context of our dark energy quantum generation process and we present a simple model. Finally, in section \\ref{phenomenology} we present the main phenomenological consequences of our model while in section \\ref{conclusions} we resume and conclude. ",
        "conclusions": "\\label{conclusions}   We will now present a summary and conclusions of our work. One of the main  goals of this paper was to understand why the dark energy is manifested at such a late time. To achieve this we have presented a novel idea,  the quantum generation of dark energy, giving  a new interpretation of the late time emergence of DE in terms of a late time quantum production of the quintessence $\\phi$ particles. We take a  $2\\leftrightarrow 2$ quantum  process between $\\phi$ and a relativistic  $\\vp$ particles. The scale where the $\\phi$ field is generated  is dynamically determined by the condition $\\Gamma/H=E_{gen}/E\\geq 1$ giving an energy scale $E\\leq E_{gen} $ with $E_{gen} = c_{gen} \\,  g^2 \\, m_{pl}$ and $c_{gen} \\simeq 6 \\cdot 10^{-4}$. Therefore the smallness of $E_{gen}$ is due to a small coupling $g$ and for $g\\simeq 10^{-12}$ gives a $E_{gen } \\simeq 1 \\, eV$ and a cross section $\\sigma_{gen}\\simeq 1 pb$. The acceleration of the universe is then due to the classical evolution of $\\phi$ and determined by the scalar potential $v(\\phi)$. We have described in section \\ref{regeneration}  a universe that initially contains no $\\phi$ particles,  i.e. $n_\\phi=\\Omega_\\phi=\\dot\\phi=v(\\phi)=0$, and once the relativistic particles $\\phi$ are produced they become  a source term for the generation of the scalar potential $v(\\phi)$. Once $v(\\phi)$ has been produced the classical equation of motion gives the evolution of $\\phi$.   We show  in section \\ref{unification} that it is possible to unify inflation and dark energy using the same quintessence field $\\phi$. To achieve the unification we required  that the potential $v(\\phi)$ has two flat regions, at high energy for inflation and low energy for dark energy. In this scenario, after inflation the field $\\phi$ decays completely and reheats the universe with standard model particles. The universe expands then in a decelerating  way dominated first by radiation and later by matter. At low energies the same interaction term that gives rise to the inflaton decay accounts for the re-generation of the $\\phi$ field giving rise to dark energy. An important difference in the quantum process between $\\phi$ and $\\vp$ at high and low energies is the value of transition rate due to the size of the $\\phi$ mass,  $m^2_\\phi(E_I)\\gg m^2_\\phi(E_o)$.   We presented  in section \\ref{unification}  a simple example on how the inflation-dark energy unification can be implemented. We used a potential $v=E_I^4(1-\\arctan[\\phi/f])$ which is flat at high and low energies. The  two parameters $E_I, f$ are determined by the density perturbations  $\\delta\\rho/\\rho$ and the value of $v_o$ at present time giving $E_I=100\\,TeV, f=10^{-39} eV$. The coupling $g$  between $\\phi$ and $\\vp$ and the coupling $h$ between $\\vp$ and the SM particles are free parameters but can be taken as $g=h^2= q \\, E_I/m_{pl}=(q/100)(E_I/100\\,TeV)10^{-12}$ giving a reheating energy $E_{RH}=(q/100)(E_I/(00\\,TeV)\\, TeV$ and a generation energy  $ E_{gen}= (q/100)^2(E_I/100\\,TeV)^2 26  \\, eV$. The cross section between $\\phi$ and $\\vp$ is $\\sigma_{gen} = g^2/32 \\pi E_{gen}^2 \\simeq  pb$ quite close to cross section of WIMP dark matter with nucleons $\\sigma_{w} \\simeq pb$ at decoupling. By fixing $g=h^2=q\\,E_I$ we have determined the coupling, which set the scales of reheating and $\\phi$ re-generation, in terms of the inflation scale $E_I$ and we can reduced the number of parameters. Of course this is not the only possible choice of $g$ and $h$.\\\\ To conclude, we have presented a general framework to produce the fundamental quintessence field $\\phi$ dynamically at low energies. The energy scale is fixed by the strength of the coupling and this offers a new interpretation of the cosmological coincidence problem:   dark energy domination starts at such small energies because of  the size of the coupling constant $g$. Finally, our approach  allows for an easy implementation of inflation and dark energy unification with the standard long periods of radiation/matter domination.     \\appendix"
    },
    "0910/0910.5756_arXiv.txt": {
        "abstract": "We present interferometric CO observations of twelve z$\\sim$2 submillimetre-faint, star-forming radio galaxies (SFRGs) which are thought to be ultraluminous infrared galaxies (ULIRGs) possibly dominated by warmer dust (T$_{\\rm dust}\\,\\simgt\\,$40\\,K) than submillimetre galaxies (SMGs) of similar luminosities.  Four other CO-observed SFRGs are included from the literature, and all observations are taken at the Plateau de Bure Interferometer (PdBI) in the compact configuration.  Ten of the sixteen SFRGs observed in CO (63\\%) are detected at $>$4$\\sigma$ with a mean inferred molecular gas mass of $\\sim$2$\\times$10$^{10}$\\,\\msun. SFRGs trend slightly above the local ULIRG L$_{\\rm FIR}$-L$^\\prime_{\\rm CO}$ relation.  Since SFRGs are about two times fainter in radio luminosity but exhibit similar CO luminosities to SMGs, this suggests SFRGs are slightly more efficient star formers than SMGs at the same redshifts.  SFRGs also have a narrow mean CO line width, 320$\\pm$80\\,\\kms.  Many SMGs have similarly narrow CO line widths, but very broad features ($\\sim$900\\,\\kms) are present in a few SMGs and are absent from SFRGs. SFRGs bridge the gap between properties of very luminous $>$5$\\times$10$^{12}$\\,\\lsun SMGs and those of local ULIRGs and are consistent with intermediate stage major mergers.  We suspect that more moderate-luminosity SMGs, not yet surveyed in CO, would show similar molecular gas properties to SFRGs.  The AGN fraction of SFRGs is consistent with SMGs and is estimated to be 0.3$\\pm$0.1, suggesting that SFRGs are observed near the peak phase of star formation activity and not in a later, post-SMG enhanced AGN phase. Excitation analysis of one SFRG is consistent with CO excitation observed in SMGs (turning over beyond \\dco).  This CO survey of SFRGs serves as a pilot project for the much more extensive survey of {\\it Herschel} and {\\sc SCUBA-2} selected sources which only partially overlap with SMGs.  Better constraints on the CO properties of a diverse high-$z$ ULIRG population are needed from ALMA to determine the evolutionary origin of extreme starbursts, and what role ULIRGs serve in catalyzing the formation of massive stellar systems in the early Universe. ",
        "introduction": "\\label{s:introduction}  Ultraluminous infrared galaxies (ULIRGs) exhibit some of the most extreme star formation rates in the Universe.  The volume density of higher redshift ULIRGs peaks at z$\\sim$2-3 \\citep{chapman05a}$-$this is also the peak epoch in the cosmic star formation rate density and volume density of active galactic nuclei \\citep[AGN; e.g.][]{fan01a,richards06a}.  Not only does this indicate a possible link between supermassive black hole growth and rapid star formation, but it also signals the most active phase in galaxy evolution and formation.  ULIRGs exhibit very intense (SFR\\simgt200\\,\\Mpy), short-lived bursts ($\\tau\\,\\sim\\,$100\\,Myr) of star formation.  The possible life cycle of a ULIRG, from star-formation dominated, dust-enshrouded galaxy, to obscured AGN and then luminous quasar \\citep[e.g.][]{sanders88b,veilleux09a}, provides a testable evolutionary sequence.  The best studied ULIRGs at high redshift are submillimetre galaxies \\citep[SMGs;][]{blain02a} which are characterised by their detection at 850\\um\\, with S$_{850}\\simgt5$\\,mJy.  While SMGs put powerful constraints on galaxy evolution theories and the environments of extreme star formation \\citep{greve05a,tacconi06a,tacconi08a,chapman05a,pope06a}, their selection is susceptible to strong temperature biasing \\citep{eales00a,blain04a}.  At the mean redshift of radio SMGs, z$\\sim$2.2, observations at 850\\um\\ sample the Rayleigh-Jeans tail of blackbody emission where the observed flux density may be approximated by S$_{850}\\propto$L$_{\\rm FIR}\\,T_{dust}^{-3.5}$.  Due to the strong dependence on dust temperature, the 850\\um\\ flux density of warm-dust ULIRGs (T$_{d}\\simgt$40\\,K) might be much lower than cooler dust specimens (T$_{d}=$20-40\\,K) thus causing the warmer-dust galaxies to evade submm detection.  This selection bias suggests that a large fraction of the z$\\sim2$ ULIRG population has not been accounted for in current work on high-z star formation.  \\citet{chapman04a} describe the first observational effort to identify warm-dust ULIRGs as a population, via the selection of submm-faint radio galaxies (SFRGs) with starburst-consistent rest-UV spectra. While they were thought to be ULIRGs by their similarities to SMGs (similar radio luminosities, optical spectra, stellar masses), without detection in the far-infrared (FIR), there was no direct evidence that their luminosities were in excess of 10$^{12}$\\,\\lsun.  Ideally, detection at shorter wavelengths in the infrared, at $\\lambda\\,\\le\\,$500\\,\\um\\, must be used to confirm a ULIRG's luminosity in the absence of submm detection; \\citet{casey09b} used 70\\,\\um\\ detection to confirm that a subset of the SFRG population contains a dominating warmer-dust component with $<T_{d}>$\\,=\\,52\\,K. While a population of warm-dust ULIRGs has been shown to exist, some fundamental questions still remain unanswered: are warm-dust ULIRGs in a post-SMG AGN heated phase?  Could they be triggered by different mechanisms than the major mergers said to give rise to cold-dust SMGs?  Investigating the molecular gas content is fundamental to the characterisation of star formation properties and gas dynamics of a galaxy population.  Molecular line transitions from carbon monoxide (CO) are a direct probe of the vast gas reservoirs that are needed to fuel high star formation rates \\citep{frayer99a,greve05a,tacconi06a,tacconi08a,chapman08a}.  The gas dynamics which are derived from these observations shed light on galaxies' evolutionary sequences by measuring how disturbed their gas reservoirs are and how long they can maintain their star formation rates with the observed fuel supply.  Recent simulations work hint that a ULIRG phase may be triggered by either major merger interactions \\citep[e.g.][]{narayanan09a} or from steady bombardment from low mass fragments \\citep[e.g.][]{dave10a}; linking observations with these different evolutionary scenarios is an essential step in understanding galaxy evolution in the early Universe.  In this paper, we present CO molecular gas observations, taken with the IRAM Plateau de Bure Interferometer (PdBI), of twelve SFRGs (and four additional SFRGs from the literature) to compare the population with SMGs and other high redshift star forming galaxies.  Section \\ref{s:observations} describes the sample selection, molecular gas observations and ancillary data, while section \\ref{s:results} presents our results, in the form of derived gas and star formation quantities of the SFRG sample.  Section \\ref{s:discussion} discusses the gas properties of the sample, compares the population to other high redshift galaxies, and hypothesizes on the role of SFRGs in a broader galaxy evolution context relative to local ULIRGs and SMGs while section \\ref{s:conclusions} concludes.  Throughout, we use a $\\Lambda$ CDM cosmology with $H_{\\rm 0} = 71$\\kms~Mpc$^{-1}$, $\\Omega_{\\rm \\Lambda}=0.73$ and $\\Omega_{\\rm m}=0.27$ \\citep{hinshaw09a}. ",
        "conclusions": "\\label{s:conclusions}  We have presented CO molecular gas observations of a sample of submillimetre-faint, star-forming radio galaxies (SFRGs).  Due to their non-detection at submillimetre wavelengths and lack of dominant AGN, these ultraluminous, \\uJy\\ radio galaxies are thought to be dominated by star formation but have warmer dust temperatures than SMGs.  Out of 16 CO-observed SFRGs (12 from this paper and 4 from the literature), 10 are detected with a mean CO luminosity of $L^\\prime_{\\rm CO[1-0]}\\,\\sim\\,$2.6$\\times$10$^{10}$\\,K\\,km\\,s$^{-1}$\\,pc$^2$, which is slightly less luminous than the CO-observed SMG sample, despite being $\\sim$2$\\times$ less luminous in the radio.  We attribute the luminosity difference to a selection bias but suggest that physical driving mechanisms might differ between the very bright ($>$10$^{13}$\\,\\lsun) and moderately bright ($\\sim$10$^{12}$\\,\\lsun) populations.  High-resolution radio imaging from MERLIN+VLA shows that the radio emission in the SFRG sample is resolved and extended with mean effective radii $\\sim$2\\,kpc, suggesting that the SFRG radio luminosities are dominated by star formation rather than AGN.  The MERLIN+VLA sizes constraints are consistent with similarly analyzed SMG MERLIN+VLA sizes.  While we note that AGN do not dominate our sample, it is possible that several of our sources have non-negligible AGN due in part to their selection as radio galaxies.  Due to limited FIR data, we use the FIR/radio correlation to derive $L_{FIR}$ and then compute extinction-free star formation rates from the FIR.  The star formation efficiencies (SFEs) of SFRGs are comparable within large uncertainties to those of SMGs and local ULIRGs, even though a few sources appear to have very high SFEs \\citep[like those in ][]{chapman08a} or very low SFEs \\citep[like those in ][]{daddi08a}. Those with perceived very high SFEs are more likely AGN dominated than super efficient; their FIR luminosities as calculated from the radio are probably overestimated.  SFRGs have narrower CO line widths than the bright subsample of SMGs at the same redshifts ($\\Delta V_{\\rm SFRG}\\,\\sim\\,$320\\,\\kms\\ and $\\Delta V_{\\rm SMG}\\,\\sim\\,$530\\,\\kms), suggesting that SFRGs might have less disturbed dynamical environments.  The line width distribution is potentially suggestive of different evolutionary stages or processes between SFRGs and SMGs; however, the observed difference with SMGs could be due to a S/N or luminosity bias.  The former would mean that intrinsically broad lines would have underestimated FWHMs due to low S/N.  The latter is due to the more thorough spectroscopic sampling of the SMG population.  SMGs have higher spectroscopic completeness and also include many objects with AGN signatures in the rest-UV/optical.  Any SFRGs which have similar AGN signatures were culled from the sample, thus eliminating some of the potentially brightest SFRGs (most of the bright SMGs which have been surveyed in CO contain optical AGN).  While less luminous SMGs (at the same luminosities of SFRGs) exist, few have been observed in CO, thus it is difficult to rule out that CO luminosity might relate to the distribution in CO with line width.  Despite selection biases, we have explored the possible physical scenarios triggering warm-dust ULIRGs in contrast to the well studied cold-dust ULIRGs.  SFRGs appear to bridge the gap between the properties of $>$10$^{13}$\\,\\lsun SMGs and $\\sim$10$^{12}$\\,\\lsun. Quantitatively, their extended radio emissions suggest sizes consistent with SMGs, implying much larger dynamical masses than local ULIRGs. We show that SFRGs have the same AGN fraction as SMGs and are therefore unlikely to represent a `post-SMG' AGN turn-on phase. Luminous SMGs have been characterised as an early infall stage during a major merger, and local ULIRGs are often described at late-type major mergers.  Here, we suggest that SFRGs (and less luminous SMGs) span the range of states during peak merger interaction."
    },
    "0910/0910.4799_arXiv.txt": {
        "abstract": "{ We estimate sensitivity coefficients to variation of the fine-structure constant $\\alpha$ and electron-to-proton mass ratio $\\mu$ for microwave $\\Lambda$-type transitions in CH molecule and for inversion-rotational transitions in partly deuterated ammonia NH$_2$D. Sensitivity coefficients for these systems are large and strongly depend on the quantum numbers of the transition. This can be used for the search for possible variation of $\\alpha$ and $\\mu$. ",
        "introduction": "Discrete microwave spectra of molecules are often used for astrophysical studies of possible variation of the fine structure constant $\\alpha=e^2/(\\hbar c)$, the electron-to-proton mass ratio $\\mu=m_\\mathrm{e}/m_\\mathrm{p}$, and the nuclear $g$-factor $g_\\mathrm{n}$. \\cite{Dar03} and \\cite{CK03} pointed out that 18 cm $\\Lambda$-doublet line of OH molecule has high sensitivity to variation of $\\alpha$ and $\\mu$. \\cite{VKB04} have shown that inversion transitions in fully deuterated ammonia $^{15}$ND$_3$ have high sensitivity to $\\mu$-variation, $Q_\\mu=5.6$. According to \\cite{FK07a}, the inversion transition in non-deuterated ammonia has a slightly smaller sensitivity, $Q_\\mu=4.5$. Note that molecular rotational lines have sensitivity $Q_\\mu=1.0$. Because of that, possible variation of constants would lead to apparent velocity offsets between $\\Lambda$-doublet OH line, or ammonia inversion line on one hand and rotational molecular lines, originated from the same gas clouds, on the other hand. This method was used by \\cite{KCL05,FK07a,MFMH08,HMM09} to establish very stringent limits on $\\alpha$- and $\\mu$-variation over cosmological timescale $\\sim 10^{10}$ years.   Recently ammonia method was applied by \\cite{LMK08} and \\cite{LML09} to dense prestellar molecular clouds in the Milky Way. These observations provide a bound of a maximum velocity offset between ammonia and other molecules at the level of $|\\Delta V| \\le 28$ m/s. This bound corresponds to $|\\Delta\\mu/\\mu| \\le 3\\times 10^{-8}$, which is two orders of magnitude more sensitive than extragalactic constraints cited above. Taken at face value the measured $\\Delta V$ shows positive shifts between the line centers of NH$_3$ and other molecules and suggests a real offset $\\Delta\\mu/\\mu= (2.2\\pm 0.4_\\mathrm{stat}\\pm 0.3_\\mathrm{sys})\\times 10^{-8}$, see \\cite{LML09}.  One of the main possible sources of the systematic errors in such observations is the Doppler noise, i.e. stochastic velocity offsets between different species caused by different spatial distributions of molecules in the gas clouds (see discussions by \\cite{KCL05} and \\cite{LRK08}). Because of that it is preferable to use lines with different sensitivity to variation of fundamental constants of the same species. \\cite{KC04} and \\cite{Koz09} showed that sensitivity coefficients for $\\Lambda$-doublet spectra of OH molecule strongly depend on quantum numbers. In this paper we focus on the CH molecule and on partly deuterated ammonia NH$_2$D. The former is similar to OH and has $\\Lambda$-doublet spectrum, which is highly sensitive to variation of $\\alpha$ and $\\mu$. For the latter the rotational and inversion degrees of freedom are strongly mixed due to the broken symmetry. This leads to a significant variation of the sensitivity of different microwave transitions to $\\mu$-variarion. Note that microwave spectra of CH and NH$_2$D from the interstellar medium were detected by several groups (see, for example, \\cite{RES74,ZT85,OBR85,LRT05,LGR08}, and references therein). ",
        "conclusions": "We have shown that there are several new microwave lines with high sensitivity to possible variation of the fundamental constants, which have been observed in the interstellar medium. Moreover, one can use the lines of the same species with different sensitivities and significantly reduce the Doppler noise."
    },
    "0910/0910.3493_arXiv.txt": {
        "abstract": "{The measurement of line broadening in cool stars is in general a difficult task. In order to detect slow rotation or weak magnetic fields, an accuracy of $1$\\,km\\,s$^{-1}$ is needed. In this regime the broadening from convective motion become important. We present an investigation of the velocity fields in early to late M-type star hydrodynamic models, and we simulate their influence on \\element[][]{FeH} molecular line shapes. The M star model parameters range between $\\log{g}$ of $3.0~-~5.0$ and effective temperatures of $2500$\\,K and $4000$\\,K.} {Our aim is to characterize the $\\teff$- and $\\log{g}$-dependence of the velocity fields and express them in terms of micro- and macro-turbulent velocities in the one dimensional sense. We present also a direct comparison between 3D hydrodynamical velocity fields and 1D turbulent velocities. The velocity fields strongly affect the line shapes of \\element[][]{FeH}, and it is our goal to give a rough estimate for the $\\log{g}$ and $\\teff$ parameter range in which 3D spectral synthesis is necessary and where 1D synthesis suffices. Eventually we want to distinguish between the velocity-broadening from convective motion and the rotational- or Zeeman-broadening in M-type stars which we are planning to measure. For the latter \\element[][]{FeH} lines are an important indicator.} {In order to calculate M-star structure models we employ the 3D radiative-hydrodynamics (RHD) code \\texttt{CO$^5$BOLD}. The spectral synthesis on these models is performed with the line synthesis code \\texttt{LINFOR3D}. We describe the 3D velocity fields in terms of a Gaussian standard deviation and project them onto the line of sight to include geometrical and limb-darkening effects. The micro- and macro-turbulent velocities are determined with the ``Curve of Growth'' method and convolution with a Gaussian velocity profile, respectively. To characterize the $\\log{g}$ and $\\teff$ dependence of \\element[][]{FeH} lines, the equivalent width, line width, and line depth are regarded.} {The velocity fields in M-stars strongly depend on  $\\log{g}$ and $\\teff$. They become stronger with decreasing $\\log{g}$ and increasing $\\teff$. The projected velocities from the 3D models agree within $\\sim~100$\\,m\\,s$^{-1}$ with the 1D micro- and macro-turbulent velocities. The \\element[][]{FeH} line quantities systematically depend on $\\log{g}$ and $\\teff$.} {The influence of hydrodynamical velocity fields on line shapes of M-type stars can well be reproduced with 1D broadening methods.  \\element[][]{FeH} lines turn out to provide a mean to measure $\\log{g}$ and $\\teff$ in M-type stars.  Since different \\element[][]{FeH} lines behave all in a similar manner, they provide an ideal measure for rotational and magnetic broadening. } ",
        "introduction": "Most of our knowledge about stars comes from spectroscopic investigation of atomic or molecular lines. In sun-like and hotter stars, the strength and shape of atomic spectral lines provides information on atmospheric structure, velocity fields, rotation, magnetic fields, etc. Measuring the effects of velocity fields on the shape of spectral lines requires a spectral resolving power between $R \\sim 10,000$ ($\\Delta v = 30$\\,km\\,s$^{-1}$) for rapid stellar rotation, $R \\ga 30,000$ ($\\Delta v = 10$\\,km\\,s$^{-1}$) for slower rotation and high turbulent velocities, and resolution on the order of $R \\sim 100,000$ for the analysis of Zeeman splitting and line shape variations due to slow convective motion.  In slowly rotating sun-like stars, usually a large number of relatively isolated spectral lines are available for the investigation of Doppler broadened spectral lines. These lines are embedded in a clearly visible continuum allowing a detailed analysis of individual lines at high precision. At cooler temperature, first the number of atomic lines is increasing so that more and more lines become blended rendering the investigation of individual lines more difficult.  At temperatures around 4000\\,K, molecular lines, predominantly \\element[][]{VO} and \\element[][]{TiO}, start to become important. At optical wavelengths, molecular bands in general consist of many lines that are blended so that the absorption mainly appears as an absorption band; individual molecular lines are difficult to identify. At temperatures in the M type stars regime (4000\\,K and less), atomic lines start to vanish because atoms are mainly neutral and higher ionization levels are weakly populated. Only alkali lines appear that are strongly affected by pressure broadening.  Thus, the detailed spectroscopic investigation of velocity fields in M dwarfs is very difficult at optical wavelengths.  M-type stars emit the bulk of their flux at infrared wavelengths redward of 1\\,$\\mu$m. This implies that observation of high SNR spectra in principle is easier in the infrared. Furthermore, M type stars exhibit a number of molecular absorption bands in the infrared, for example \\element[][]{FeH}.  In these bands, the individual lines are relatively well separated and provide a good tracer of stellar velocity fields. The lines are intrinsically much narrower than atomic lines in sun-like stars because Doppler broadening due to the temperature related motion of the atoms and molecules is much reduced. Thus, the lines can be used for the whole arsenal of line profile analysis that has been applied successfully to sun-like stars over the last decades.  Examples of analyses using \\element[][]{FeH} lines are the investigation of the rotation activity connection in field M-dwarfs, which requires the measurement of rotational line broadening with an accuracy of $1$\\,km\\,s$^{-1}$ \\citep{2007A&A...467..259R}. Another example is the measurement of magnetic fields comparing Zeeman broadening in magnetically sensitive and insensitive absorption lines \\citep[see e.g.~][]{2006ApJ...644..497R}. A precise analysis of \\element[][]{FeH} lines, however is only possible if the underlying velocity fields of the M dwarfs atmospheres are thoroughly understood. In this paper, we model the surface velocity fields of M type stars and their influence on the narrow spectral lines of \\element[][]{FeH}.  We calculate 3D-\\texttt{CO$^5$BOLD} structure models \\citep{2002A&A...395...99L} which serve as an input for the line formation program \\texttt{LINFOR3D} \\citep[based on][]{Bascheck1966}. Turbulence's are included in a natural way using hydrodynamics, so that we are able to investigate the modeled spectral lines for effects from micro- and macro-turbulent velocities in the classical sense and their influence on the line shapes. The comparison with 1D-models gives a rough estimate of the necessity of using 3D-models in the spectral domain of cool stars. In the first part of this paper we investigate the velocity fields in the models and their dependence on $\\log{g}$ and $\\teff$. In the second part, we investigate the influence of velocity fields, $\\log{g}$, and $\\teff$ on the \\element[][]{FeH} molecular lines. ",
        "conclusions": "We investigated a set of M-star models with $\\teff=2500$\\,K - $4000$\\,K and $\\log{g}=3.0$ -- $5.0$ [cgs]. For these models, the 3D hydrodynamic radiative transfer code \\texttt{CO$^5$BOLD} was used. The horizontal and vertical velocity fields in the 3D models were described with a binning method. The convective turn-over point is clearly visible in the atmospheric velocity dispersion structure. To investigate the influence of these velocity fields on spectral line shapes, a description for geometrical projection and limb-darkening effects was applied. With the use of contribution functions, we took only these parts in the atmosphere into account where the lines were formed. The resulting velocity dispersions range from $400$ -- $1600$\\,m\\,s$^{-1}$ with decreasing $\\log{g}$ and with increasing $\\teff$ from $200$ -- $1400$\\,m\\,s$^{-1}$. These values agree well with velocities deduced from line shapes. We expressed the hydrodynamical velocity fields of the 3D models in terms of the classical micro- and macro-turbulent velocities. With this description and the obtained micro- and macro-turbulent velocities, it is possible to reproduce 3D spectral lines on 1D atmosphere models very accurately, hence time consuming 3D treatment of \\element[][]{FeH} molecular lines in the regime of cool stars is not necessary for line profile analysis. A comparison of our velocities with a set of velocities determined from observations with spectral fitting methods showed that the macro-turbulent velocities agree, but the micro-turbulent velocities are a factor of two or three smaller than the ones determined from observations.  A line shift due to the larger up-flowing area in the convection zone was investigated too. It is on the order of a few m/s up to $50$\\,m\\,s$^{-1}$ for a very low gravity model. The time dependent jitter in line positions is only about m/s and would be reduced to \\,mm\\,s$^{-1}$ in a real star, due to high number of contributing elements.  In order to use \\element[][]{FeH} molecular lines for investigations of spectroscopic/physical properties in cool stars (e.g. Zeeman- or rotational broadening), we explored the behavior in a set of lines on $\\log{g}$ and $\\teff$. We investigated ten \\element[][]{FeH} lines between 9950\\AA\\ and 9990\\AA\\ on our models with the spectral synthesis code \\texttt{LINFOR3D}. \\element[][]{FeH} lines react on different effective temperatures as expected due to the change in chemical composition and pressure. The lines showed also a weak dependence on surface gravity due to changing densities and pressure. The broadening from velocity fields in the 3D models of the $\\log{g}$ series is very strong, but for the $\\teff$ series the broadening from velocity fields is almost covered by v.d.Waals broadening. The difference in line width for hot models is up to $0.5$\\,km\\,s$^{-1}$ and for low gravity models around $1$\\,km\\,s$^{-1}$.  That means for the 1D spectral synthesis, that one has to include correct micro- and macro-turbulent velocities for small surface gravities or hot $\\teff$. Due to the fact that the FWHM $\\log{g}$ dependence of \\element[][]{FeH} lines goes in the opposite direction as the $\\log{g}$ dependence of the velocity fields, the \\element[][]{FeH} lines become a great mean to measure surface gravities in cool stars. Because the velocity fields scale with $\\log{g}$ and it should be easily possible to detect them.  \\element[][]{FeH} lines with different quantum numbers do not show significant differences for both, $\\log{g}$- and $\\teff$-series. That means the broadening of the lines does not depend on $J$, $\\Omega$, or the branch. Furthermore lines with weak magnetic sensitivity behave just like lines with strong magnetic sensitivity. All lines are broadened in the same way by thermal and hydrodynamical motions. Only the transition probability expressed in the $\\log{gf}$ value influences the behavior of the lines. The line with the lowest $gf$-values did not saturate at low $\\teff$, but in general they are similar to the other \\element[][]{FeH} lines.  It is possible to treat the \\element[][]{FeH} molecular lines with different quantum numbers as homogenous in the absence of magnetic fields. That allows to use \\element[][]{FeH} lines to measure magnetic fields \\citep{2006ApJ...644..497R,2007ApJ...656.1121R}. Hence we conclude, that these lines also are an appropriate mean to measure magnetic field strength in M-type stars."
    },
    "0910/0910.3941_arXiv.txt": {
        "abstract": "We study for what specific values of the theoretical parameters the axion can form the totality of cold dark matter. We examine the allowed axion parameter region in the light of recent data collected by the WMAP5 mission plus baryon acoustic oscillations and supernovae \\cite{komatsu}, and assume an inflationary scenario and standard cosmology. We also upgrade the treatment of anharmonicities in the axion potential, which we find important in certain cases. If the Peccei-Quinn symmetry is restored after inflation, we recover the usual relation between axion mass and density, so that an axion mass $m_a =(85\\pm3){\\rm~\\mu eV}$ makes the axion 100$\\%$ of the cold dark matter. If the Peccei-Quinn symmetry is broken during inflation, the axion can instead be 100$\\%$ of the cold dark matter for $m_a < 15{\\rm~meV}$ provided  a specific value of the initial misalignment angle $\\theta_i$ is chosen in correspondence to a given value of its mass $m_a$. Large values of the Peccei-Quinn symmetry breaking scale correspond to small, perhaps uncomfortably small, values of the initial misalignment angle $\\theta_i$. ",
        "introduction": "About 84$\\%$ of the non-relativistic matter in the Universe is in the form of cold dark matter (CDM)\\cite{komatsu}, whose nature is still to be discovered. One of the most promising hypothetical candidates that could account for the CDM observed is the axion \\cite{weinberg, peccei}.  The properties of the axion as the CDM particle have been studied in various papers \\cite{preskill, beltran}. We examine the possibility that the axion accounts for 100$\\%$ CDM in the light of the WMAP5 mission, baryon acoustic oscillations (BAO) and supernovae (SN) data. We also upgrade the treatment of anharmonicities in the axion potential, which we find important in a specific region of the parameter space. The axion parameter space is described by three quantities, the PQ energy scale $f_a$, the Hubble parameter at the end of inflation $H_I$ and the axion initial misalignment angle $\\theta_i$. ",
        "conclusions": ""
    },
    "0910/0910.3036_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\par\\noindent The cores of Active Galactic Nuclei (AGN), identified as quasars, emit a huge amount of power at visible and ultraviolet frequencies \\cite{bmp}. It obtains its power by the gravitational potential energy of a massive black hole residing at its center \\cite{kori}. The radiation is emitted by the accretion disk surrounding the black hole. In this paper we compute the luminosity of the invisible axions \\cite{PQ1,PQ2,Weinberg,Wilczek,McKay,Kim1,Dine,Zhitnitsky,McKay2,Kim} from AGNs. We also consider a hypothetical light pseudoscalar whose couplings and mass are not related to one another. Our motivation for this study is two folds. The pseudoscalar flux from AGNs may be used to impose limits on its mass and couplings. If the pseudoscalar flux is sufficiently large then it might also provide an explanation for the observed large scale alignment of visible polarizations from quasars. Large scale alignment, on distance scales of a Gpc, has been observed in many regions of the sky \\cite{huts98,huts01,huts02,JNS04}. A statistically significant signal of alignment with the local supercluster has also been observed \\cite{huts98,huts01,huts02,JNS04,JP04}. This effect may be explained in terms of the conversion of photons to pseudoscalars in the local supercluster magnetic field. However this explanation is not consistent with data. The problem arises due to the observed difference in the distribution of polarizations among the Radio Quiet (RQ) and optically selected (O) quasars and the Broad Absorption Line (BAL) quasars. The polarization distribution of the RQ and O quasars peaks at very low values. The magnitude of the mixing required to explain the alignment effect is sufficiently large so as to completely wash out this difference. In Ref. \\cite{JPS02} it was suggested that if the pseudoscalar flux from quasars is sufficiently large at visible frequencies, than conversion in the supercluster magnetic field may consistently explain the alignment with supercluster. In this case the alignment is explained in terms of the conversion of pseudoscalars to photons.  We first study the emission of pseudoscalars from the accretion disk via the Compton, Bremsstrahlung and the Primakoff channels. In this calculation we assume the pseudoscalar to be the standard axion. Besides emission from the accretion disk, pseudoscalars may also be produced in the AGN atmospheres due to the conversion of photons to pseudoscalar in the background magnetic field. The probability for this conversion is negligible for the standard axion but can be large if the pseudoscalar mass is very small. ",
        "conclusions": "\\par\\noindent In this paper, we have found that the luminosity of pseudoscalars from the AGN accretion disk due to Compton, Bremsstrahlung and  Primakoff channels is very small in comparison to the photon luminosity. However, the photons in visible and ultraviolet frequencies may convert to pseudoscalars outside the accretion disk due to pseudoscalar-photon mixing in the background magnetic field. Taking extinction of photons into account and using the current limit on the pseudoscalar-photon coupling, we find that the pseudoscalar flux produced by this process is relatively large. For ultraviolet frequencies, which would be observed in the visible range on earth, this flux may dominate the  photon flux. A large pseudoscalar flux may provide a consistent explanation for the large scale coherent orientation of the visible polarizations from quasars. \\medskip \\\\"
    },
    "0910/0910.2426_arXiv.txt": {
        "abstract": "We present results from a multi-month reverberation mapping campaign undertaken primarily at MDM Observatory with supporting observations from around the world. We measure broad line region (BLR) radii and black hole masses for six objects.  A velocity-resolved analysis of the H$\\beta$ response shows the presence of diverse kinematic signatures in the BLR. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.1878_arXiv.txt": {
        "abstract": "The effect of the chosen analysis energy window on the results of a dark matter experiment is exemplified by the curious intersection of the exclusion plots of the XENON10 and the CDMS experiments. After proving that the narrow energy window XENON10 chose to analyze is indeed the cause of such intersection, a method to determine the high-energy extreme of the recoil energy window an experiment should use is obtained. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.1094_arXiv.txt": {
        "abstract": "We use new large area far infrared maps ranging from 65-500\\,$\\mu$m obtained with the AKARI and the Balloon-borne Large Aperture Submillimeter Telescope (BLAST) missions to characterize the dust emission toward the Cassiopeia~A supernova remnant (SNR). Using the AKARI high resolution data we find a new ``\\emph{tepid}'' dust grain population at a temperature of $\\sim35$\\,K and with an estimated mass of 0.06\\,M$_{\\odot}$.  This component is confined to the central area of the SNR and may represent newly-formed dust in the unshocked supernova ejecta.  While the mass of tepid dust that we measure is insufficient by itself to account for the dust observed at high redshift, it does constitute an additional dust population to contribute to those previously reported.  We fit our maps at 65, 90, 140, 250, 350, and 500\\,$\\mu$m to obtain maps of the column density and temperature of ``cold'' dust (near 16~K) distributed throughout the region. The large column density of cold dust associated with clouds seen in molecular emission extends continuously from the surrounding interstellar medium to project on the SNR, where the foreground component of the clouds is also detectable through optical, X-ray, and molecular extinction.  At the resolution available here, there is no morphological signature to isolate any cold dust associated only with the SNR from this confusing interstellar emission. Our fit also recovers the previously detected ``hot'' dust in the remnant, with characteristic temperature 100~K. ",
        "introduction": "\\label{The_CasA_SNR} Determining the origin of cosmic dust is fundamental to our understanding of many astronomical processes, including star formation and galaxy evolution. Galaxies and quasars at high redshift have been found to contain large amounts of dust \\citep[$\\ge10^8$\\,M\\subsun;][]{Dunlop_1994, Archibald_dust_in_distant_quasars_2001, Isaak_2002, Priddey_2008}, at a time when the Universe was only about one tenth of its present age. The main source of dust \\textit{injection} within our Galaxy is thought to be the stellar winds of stars on the asymptotic giant branch (AGB) of the Hertzsprung-Russell (HR) diagram \\citep{morgan_edmunds_2003}. Stars at this early epoch would not have been able to reach the AGB phase in the available time, and therefore cannot be the source of the observed dust. Heavy elements are produced in the explosions of supernovae (SNe) and, for several years, models have predicted that considerable amounts of fresh dust (0.1-1\\,M\\subsun) could also be produced \\citep{Kozasa_1991_SNR_dust_model, Woosley_1995_SNR_dust_model, Clayton_1999_SNR_dust_model, Todini_2001_SNR_dust_model}. The life cycle of high mass stars ($>$8\\,M\\subsun), the progenitors of type II SNe, is sufficiently short for SNe to occur within the required timescales. As a result, SNe have been proposed as a possible solution for the origin of the dust seen at high-redshift. However, in order for SNe to generate sufficient dust mass to fill this gap, each supernova would need to generate 0.4-1\\,M\\subsun~dust \\citep{Dwek_2007}. This quantification does not account for dust grain destruction within the supernova, thereby making it a lower bound.  In seeking evidence regarding this hypothesis, the focus has been on dust detectable in SNe and supernova remnants (SNR), as reviewed briefly below.  However, it seems less well appreciated that injection is only part of the story.  \\citet{Draine_2003} reviews the often-ignored arguments that, at least in our Galaxy, the interstellar dust is continually processed on a timescale $3 \\times 10^8$~yr and most of its mass is formed in the interstellar medium. Nevertheless, injected dust is critical at least as seeds for further evolution of the dust population.  Studies of SNe find only trace amounts of dust in the hot ejecta, with typical masses of order 10$^{-4}$\\,M$_{\\odot}$ \\citep{Dwek_1987,Lagage_1996,Arendt_1999,Ercolano_2007,Meikle_2007}.  Dust studies in SNR appear more promising, a prime example being Cassiopeia A (Cas~A). Cas~A is the remnant of a type IIb supernova event which occurred around AD 1680 \\citep{Raymond_1984, Thorstensen_2001, Fesen_2006a, Krause_2008_TypeIIb}. The progenitor star is believed to have had a mass greater than 20\\,M\\subsun~\\citep{Perez-Rendon_2002}, and the remnant is at a distance $D \\sim 3.4$\\,kpc \\citep{Reed_1995}. Early observations of Cas~A made with IRAS \\citep{IRAS_1984} and ISO \\citep{ISO_1996} did not extend to longer wavelengths, and therefore detected only the ``hot'' ($\\sim$100\\,K) dust component, whose mass seemed insufficient to provide the levels of dust seen in high redshift galaxies. The low angular resolution also made the study of sub-structure in the remnant at intermediate wavelengths difficult. The \\textit{Spitzer Space Telescope} \\citep{Spitzer_06} has been used to study Cas~A \\citep{Hines2004,Krause_2004_CasA,ennis2006}. Most recently \\cite{Rho_2008}, exploiting the angular and spectral resolution achieved with the \\textit{Spitzer} infrared spectrograph \\citep[IRS;][]{Houck_2004}, find between 0.020 -- 0.054\\,M$_{\\odot}$ of hot dust.  This is an order of magnitude greater than that previously measured and they conclude that, within modeling uncertainties for galaxy evolution, this could be sufficient to explain at least the lower limit to the dust levels in high-redshift galaxies presented by \\cite{Isaak_2002}.  Much more controversial is the question of a ``cold'' dust component in the SNR, because of the issue of contamination by line-of-sight interstellar emission from dust which is also cold ($\\sim$16\\,K; \\S~\\ref{comp_sep}). \\cite{Dunne_2003_CasA}, using data from the Submillimetre Common User Bolometer Array \\citep[SCUBA;][]{Holland_SCUBA_99}, find evidence for $2-4\\,{\\rm M\\subsun}$ of dust in Cas~A at a temperature of $\\sim$18\\,K, significantly more than the mass of hot dust. Using the same methodology, \\citet{Morgan_2003_Kepler} find $\\sim$1\\,M$_{\\odot}$ of cold dust in Kepler's supernova remnant, a thousand times greater than previous measurements for this SNR. Both of these remnants are sufficiently young for the dust observed to be freshly formed in the remnant, rather than being material swept up from the ISM by the shock-wave \\citep{Dickel_1988, Hughes_1999, Synch_Wright_1999}. More recently, \\cite{Dunne_2009} presented further evidence for cold dust in Cas~A using SCUBA polarization data. These show dust emission polarized in an orientation consistent with that of the magnetic field deduced from the radio synchrotron emission, suggesting the detected dust is in the SNR. They find a conservative lower limit for this dust mass of 1\\,M$_{\\odot}$.  They also attribute apparent depolarization at the brightest feature to dilution by line of sight interstellar emission.  \\cite{Dwek_2004} has argued that the mass estimates in \\cite{Dunne_2003_CasA} exceed the total anticipated mass ejection of Cas~A and suggest that if the submillimeter observations are valid they might instead imply the presence of a much smaller amount of dust which is a much more efficient millimeter wavelength radiator, such as iron needles.  The alignment of such needles could affect the polarization.  By correlating cold dust emission with molecular line absorption against the SNR synchrotron emission, \\cite{Krause_2004_CasA} argued that the dust is, in fact, associated with a molecular cloud located along the line of sight to the SNR.  Molecular emission has also been studied in this direction (e.g., \\citealt{liszt_lucas_1999}), showing that the cloud(s) extend well beyond the SNR itself. \\citeauthor{Krause_2004_CasA} estimate the fresh dust yield within Cas~A to be at least an order of magnitude lower than that found by \\cite{Dunne_2003_CasA}. Even so, this would still provide the predominant dust mass in Cas~A, bolstering the possibility of explaining the quantities of dust seen at high-redshift.  The emission from cold dust peaks in the far-infrared and submillimeter and so neither SCUBA, on the long wavelength side, nor Spitzer, on the short side, is ideal for isolating a cold dust component.  Near the thermal peak, the best large scale maps covering Cas~A and its environs are from the Multiband Imaging Photometer for Spitzer \\citep[MIPS;][]{MIPS_2004} at 160\\,$\\mu$m \\citep{Krause_2004_CasA} and from the ISO Serendipity Survey at 170\\,$\\mu$m \\citep{Stickel_2007}.  Neither map has full coverage, with MIPS missing data on small scales and ISOSS on larger.  We observed Cas~A using the Far-Infrared Surveyor \\citep[FIS;][]{Kawada_2007} instrument on-board AKARI. We obtained fully-sampled images of sub-arcminute resolution covering a wide area surrounding Cas~A in four photometric bands from 50 to 180\\,$\\mu$m. This is the wavelength range over which the emission from newly formed hot dust becomes faint and emission from cold dust begins to dominate. Therefore, these AKARI FIS images are very useful for investigating the presence of colder components of dust in the remnant.  The Balloon-borne Large-Aperture Submillimeter Telescope\\footnote{\\tt www.blastexperiment.info} \\citep[BLAST;][]{Pascale_2008} was also used to observe the Cas~A region at 250, 350, and 500\\,$\\mu$m. These bands fill in the cold dust spectral energy distribution (SED) on the long wavelength side of the peak. The high mapping speed of BLAST means it was possible to cover a large area surrounding the SNR, giving data for both the SNR and the interstellar cloud structure in the surrounding region.  By contrast, the prior SCUBA maps are of a small area and involve deconvolution of a three-beam chopping pattern which reconstructs large scale power poorly. AKARI and BLAST have the advantage that, like \\textit{Spitzer}, they are not required to perform chopped observations. As a result, our new maps are sensitive to the large scale structure present in the Cas~A field. The relatively large area of these maps, coupled with their sensitivity and wavelength coverage, make them ideal for investigation of cold dust emission from Cas~A and the interstellar clouds.  We describe the AKARI and BLAST observations in \\S~\\ref{data}. Several distinct sources of emission are distinguishable in these maps. Taking advantage of the AKARI spatial resolution, in \\S~\\ref{internal_dust} we identify a new morphologically compact source of ``tepid'' ($\\sim$35\\,K) dust emission centered on the SNR whose spectrum is also distinct, peaking between the SED peak of the hot dust in the SNR and that of the cold dust. In \\S~\\ref{photometry} we perform aperture photometry on the maps in the region of the Cas~A SNR.  The resulting global SED further illustrates the different spectral components and shows that without the additional morphological information it is not possible to unambiguously distinguish the tepid dust component. In \\S~\\ref{comp_sep} we fit the six-band AKARI-BLAST data with a simple spectral model to make column density and temperature maps for the cold dust.  These clearly illustrate the confused nature of cold dust emission on the line of sight to the SNR. In \\S~\\ref{foreground} the derived column density on the line of sight is compared to that obtained by other techniques. We present our conclusions in \\S~\\ref{conclusions}. ",
        "conclusions": "\\begin{enumerate} \\item{We presented far-infrared/submillimeter data at 65 -- 500\\,$\\mu$m for the Cas~A supernova remnant and the surrounding region. We used these maps to characterize the interstellar dust emission using data from cloud regions well beyond the SNR.} \\item{We used high resolution ARAKI data to probe the spectral region between the hot dust emission from the SNR shock-front and cold dust emission. Using a spectrum-informed clean technique we identified a new tepid dust population at temperature of $\\sim$35\\,K. The mass of this individual dust population was estimated to be 0.06\\,M$\\odot$, but with considerable uncertainty because of its dependence on the choice of $\\kappa$.} \\item{The dust yield for this new and independent tepid component is comparable to that estimated previously for the hot dust component by \\cite{Rho_2008}. While such yields could contribute to the dust masses seen in high redshift galaxies, they are still less than the required level 0.4-1\\,M\\subsun estimated by \\citet{Dwek_2007}. While the mass we measure is insufficient to account for the dust observed at high redshift, when taken in combination with the hot and cold dust masses previously reported by \\cite{Rho_2008} and \\cite{Dunne_2009}, it strengthens the argument for supernovae as a potentially significant source of dust production in the high-redshift universe.} \\item{We developed a simple physically-motivated model of the SED of the SNR and interstellar emission and fit this to six-wavelength bands at each pixel. From this we obtained temperature and dust mass column density maps.  The interstellar dust was found to be at a temperature of $\\sim 16.5$\\,K, in keeping with previous measurements, but now better constrained due to the improved wavelength coverage.} \\item{We show that the high level of confusion arising from the   interstellar cloud structure projected on the SNR precludes a   significant detection of cold dust directly associated with Cas~A.   The same source of confusion will have affected previous estimates   of cold dust in Cas~A, increasing the uncertainty of those   estimates.  This analysis was not sufficiently sensitive to identify the lower limiting mass found by \\cite{Dunne_2009} or the lower value by \\cite{Krause_2004_CasA} and therefore does not preclude the possibility of a significant population of cold SNR dust grains with temperature close to that of the interstellar dust. The higher angular resolution data anticipated with \\textit{Herschel} working close to the peak of the cold dust emission, together with correlations with surrogates of the interstellar column density, could result in a more sensitive probe.} \\end{enumerate}"
    },
    "0910/0910.3869_arXiv.txt": {
        "abstract": "We present infrared interferometric imaging of the S-type Mira star $\\chi$ Cygni. The object was observed at four different epochs in 2005-2006 with the IOTA optical interferometer (H\\ band). Images show up to $40\\%$ variation in the stellar diameter, as well as significant changes in the limb darkening and stellar inhomogeneities. Model fitting gave precise time-dependent values of the stellar diameter, and reveals presence and displacement of a warm molecular layer.  The star radius, corrected for limb darkening, has a mean value of $12.1\\,$mas and shows a $5.1\\,$mas amplitude pulsation.  Minimum diameter was observed at phase $0.94\\pm0.01$. Maximum temperature was observed several days later at phase $1.02\\pm0.02$.  We also show that combining the angular acceleration of the molecular layer with CO ($\\Delta v = 3$) radial velocity measurements yields a $5.9\\pm1.5\\,$ mas parallax. The constant acceleration of the CO molecules -- during 80\\% of the pulsation cycle -- lead us to argument for a free-falling layer. The acceleration is compatible with a gravitational field produced by a $2.1^{+1.5}_{-0.7}$ solar mass star. This last value is in agreement with fundamental mode pulsator models.  We foresee increased development of techniques consisting in combining radial velocity with interferometric angular measurements, ultimately allowing total mapping of the speed, density, and position of the diverse species in pulsation driven atmospheres. ",
        "introduction": "Mira variables are low to intermediate mass AGB stars that pulsate with a period of about 1 year. They have a cool ($T_{\\rm eff}\\leq 3000\\,K$) and extended ($R>100\\,R_\\sun$) photosphere. As such, they are bright ($M_k \\leq -7$) infrared beacons, individually observable far into galaxies of the Local Group \\citep{1999IAUS..191..551Z}. They have the potential to probe places where the distance \\citep[eg, NGC\\,5128 in][]{2004A&A...413..903R} or reddening \\citep[eg, the Galactic Center in][]{2009MNRAS.tmp.1192M} does not allow observation of the fainter/bluer -- and rarer -- Cepheids.  However, the challenge to overcome is that Mira stars are both intrinsically complicated and ill-understood.  Two important relations are of special interest: the period/luminosity (P/L) and the period/mass/radius (P/M/R).  The first relation has been derived from population studies (sequence ``C'' in the LMC from \\citet{2000PASA...17...18W} and also in the globular cluster 47 Tuc from \\citet{2005A&A...441.1117L}).  The present best parameterization of the P-L relation within our galaxy is \\citep{2008MNRAS.386..313W}: \\begin{equation} M_k=-(3.51\\pm0.20)(\\log(P)-2.38)-(7.25\\pm0.07)\\,, \\label{eq:PL} \\end{equation} where $P$ is the period in days.  The zero point of this relation is the most uncertain parameter, with its dependence on the metallicity hardly known. The main difficulty is that parallax values are inaccurate and error-prone due to the large size and inhomogeneous surface brightness of the objects.  The second relation, the P/M/R relation, has more relevance to the fundamental physics of the star. It is extremely dependent on the pulsation mode, but also, less crucially, on the surface density and metallicity. The P/M/R relation has been formally derived from numerical modeling of the fundamental pulsation mode of theses stars \\citep{Wood..89}: \\begin{equation} \\log(P)=-2.07+1.94 \\log(R/R_\\sun)-0.9 \\log(M/M_\\sun)\\,, \\label{eq:PMR} \\end{equation} Twenty years later, this model-derived relationship has still rarely been confronted with observation. This paper is a first step forward to establish the P/L and P/M/R relations on a new firm observational footing.  Of the three crucial parameters (distance, mass, and radius), the angular diameter is far from being the easiest value to obtain. Because the surface gravity is several orders of magnitude lower than the sun, the pulsation of Mira variable leads to an extended atmosphere. In the cool upper layers, significant amount of the products of the helium fusion react to form di- and polyatomic molecules including TiO, SiO, CO and H$_2$O.  The forest of molecular lines and scattering from the dust lead to exotic intensity distributions not at all like a simple stellar disk. In the past, this substantially affected many stellar angular diameters measurements \\citep{1996AJ....112.2147V,1999A&A...345..221P}, paving the way to contentious discussion on the mode of pulsation \\citep{1998A&A...333..647B,1999ApJ...514L..35Y}.   However, nowadays, interferometers are able to provide maps of the brightness distribution as a function of wavelength \\citep{2009A&A...496L...1L,2009ApJ...700..114P}. Images of the Mira star T\\ Lep revealed a shell-like atmosphere, with a bright chromatic zone distinctly detached from the photosphere. This could be the first image of what \\citet{2004A&A...424.1011O} called the MOLsphere, a zone of increased density in which formation of warm molecular species would be favored. Accounting for this layer is the key to obtain a correct value for the diameter \\citep{1999A&A...345..221P,2002ApJ...579..446M,2004A&A...426..279P}. We will also show that we can apply to this layer a modified Baade-Wesselink method to derive the distance and mass of the star.   The test star of this paper is $\\chi$ Cyg, a S-type Mira star. It has a pulsation period of 408 days, a photometric magnitude ranging from 5.3 to 13.3, and intense emission lines at postmaximum \\citep{1947ApJ...106..274M}. This suggests a large pulsation amplitude.  Images were obtained with the IOTA interferometer at four different stellar phases, chronologically $\\phi = 0.93$, 0.26, 0.69 and 0.79. In the next section, we describe the observations and give a short overview of the dataset. In section~\\ref{sc:imaging}, we use an image reconstruction algorithm to map the brightness distribution of the star. Precise geometrical parameters of the star, including existence of the molecular layer, are determined by model fitting in section~\\ref{sc:param_image}. From these values, temperatures and opacities are deduced in section~\\ref{sec:Physics}. Finally, in section~\\ref{sc:der_ma}, we combine angular acceleration with radial velocities measurements to derive estimations of the distance and mass of the star. ",
        "conclusions": ""
    },
    "0910/0910.0437_arXiv.txt": {
        "abstract": "We performed an optical spectroscopic monitoring of the blazar 3C 454.3 from September 2003 to July 2008. Sixteen optical spectra were obtained during different runs, which constitute the first spectroscopic monitoring done in the rest-frame UV region (z=0.859). An overall flux variation of the Mg{\\small II}($\\lambda$2800 \\AA) by a factor $\\sim3$ was observed, while the corresponding UV continuum ($F_{\\rm cont}$ at $\\lambda$3000 \\AA) changed by a factor $\\sim 14$. The Mg{\\small II} emission lines respond proportionally to the continuum variations when the source is in a low-activity state. In contrast, near the optical outbursts detected in 2005 and 2007, the Mg{\\small II} emission lines showed little response to the continuum flux variations. During the monitored period the UV Fe{\\small II} flux changed by a factor $\\sim 6$ and correlated with $F_{\\rm cont}$ ($r=0.92$). A negative correlation between EW(Mg {\\small II}) and $F_{\\rm cont}$ was found, i.e.\\ the so-called ``Intrinsic Baldwin Effect''. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.2118_arXiv.txt": {
        "abstract": "The age and chemical composition of the stars in present-day galaxies carry important clues about their star formation processes. The latest generation of population synthesis models have allowed to derive age and stellar metallicity estimates for large samples of low-redshift galaxies. After reviewing the main results about the distribution in ages and metallicities as a function of galaxy mass, I will concentrate on recent analysis that aims at disentangling the dependences of stellar populations properties on environment and on galaxy stellar mass. Finally, new models that predict the response of the full spectrum to variations in [$\\alpha$/Fe] will allow us to derive accurate estimates of element abundance ratios and gain deeper insight into the timescales of star formation cessation. ",
        "introduction": "The age and chemical composition of the stellar populations in galaxies, together with galaxy mass, are key ingredients to uncover galaxy formation and evolutionary paths. Estimates of stellar populations parameters are derived by interpreting detailed spectral information, such as absorption features, on the basis of stellar population synthesis (SPS) models. Our ability of interpreting galaxy spectra has greatly improved with the development of SPS models that i) predict the full spectrum of simple stellar populations (SSPs) at medium/high resolution, allowing to adjust the models to the data quality rather than the other way round, and ii) have a better coverage of the stellar parameters space, allowing the interpretation of a broader range of stellar populations \\citep{1999ApJ...513..224V,2003MNRAS.344.1000B,2007MNRAS.382..498C}.  Together with the development of SPS models and spectral fitting techniques, the statistical power of large spectroscopic surveys has allowed to put on a firm ground our understanding of the stellar populations in nearby galaxies. In this contribution I will briefly review results on the dependence of stellar populations properties on galaxy mass, as obtained from the analysis of Sloan Digital Sky Survey (SDSS) galaxy spectra. I will then discuss to what extent these relations are shaped by the environment in which galaxies reside, as recently analysed in Pasquali et al (2009, submitted) combining the SDSS DR4 catalogue of stellar ages and metallicities with the \\cite{2007ApJ...671..153Y} group catalogue.  I will conclude with a brief outlook on new population synthesis models that allow accurate estimates of $\\alpha$-element abundance ratios. ",
        "conclusions": ""
    },
    "0910/0910.3214_arXiv.txt": {
        "abstract": "We investigate the quality of associations of astronomical sources from multi-wavelength observations using simulated detections that are realistic in terms of their astrometric accuracy, small-scale clustering properties and selection functions. We present a general method to build such mock catalogs for studying associations, and compare the statistics of cross-identifications based on angular separation and Bayesian probability criteria. In particular, we focus on the highly relevant problem of cross-correlating the ultraviolet GALEX and optical SDSS surveys. Using refined simulations of the relevant catalogs, we find that the probability thresholds yield lower contamination of false associations, and are more efficient than angular separation. Our study presents a set of recommended criteria to construct reliable crossmatch catalogs between SDSS and GALEX with minimal artifacts. ",
        "introduction": "Astrophysical studies can gain significantly by associating data from different wavelength ranges of the electromagnetic spectrum. Dedicated multi-wavelength surveys have been a strong focus of observational astronomy in recent years, e.g. AEGIS \\citep{Davis_2007}, COSMOS \\citep{Scoville_2007}, or GOODS \\citep{Dickinson_2003}. At redshifts lower than those probed by these surveys, several surveys of NASA's Galaxy Evolution Explorer \\citep[GALEX;][]{Martin_2005} essentially provide the perfect ultraviolet counterparts of the Sloan Digital Sky Survey \\citep[SDSS;][]{York_2000} optical data sets. These surveys or the combination of these datasets enables to provide invaluable insights on stars and galaxy properties.     Naturally, these data are taken by different detectors of the separate projects, hence it is required to combine their information by associating the independent detections. Recent work by \\citet{Budavari_2008} laid down the statistical foundation of the cross-identification problem. Their probabilistic approach assigns an objective Bayesian evidence and subsequently a posterior probability to each potential association, and can even consider physical information, such as priors on the spectral energy distribution or redshift, in addition to the positions on celestial sphere. In this paper, we put the Bayesian formalism to work, and aim to assess the benefit of using posterior probabilities over simple angular separation cuts using mock catalogs of GALEX and SDSS. In Section~\\ref{sec_simulations}, we present a general procedure to build mock catalogs that take into account source confusion and selection functions. Section~\\ref{sec_xmatch} provides the details of the cross-identification strategy, and defines the relevant quality measures of the associations based on angular separation and posterior probability. In Section~\\ref{sec_results}, we present the results for the GALEX-SDSS cross-identification, and propose a set of criteria to build reliable combined catalogs. ",
        "conclusions": ""
    },
    "0910/0910.4328_arXiv.txt": {
        "abstract": "We made mid-infrared observations of the 10\\,M$_\\odot$ Herbig Be star HD200775 with the Cooled Mid-Infrared Camera and Spectrometer (COMICS) on the 8.2\\,m Subaru Telescope. We discovered diffuse emission of an elliptical shape extended in the north-south direction in $\\sim$1000\\,AU radius around unresolved excess emission. The diffuse emission is perpendicular to the cavity wall formed by the past outflow activity and is parallel to the projected major axis of the central close binary orbit. The centers of the ellipse contours of the diffuse emission are shifted from the stellar position and the amount of the shift increases as the contour brightness level decreases. The diffuse emission is well explained in all of geometry (the shape and the shift), size, and configuration by an inclined flared disk where only its surface emits the mid-infrared photons. Our results give the first well-resolved infrared disk images around a massive star and strongly support that HD200775 is formed through the disk accretion. The disk survives the main accretion phase and shows a structure similar to that around lower-mass stars with 'disk atmosphere'. At the same time, the disk also shows properties characteristic to massive stars such as photoevaporation traced by the 3.4\\,mm free-free emission and unusual silicate emission with a peak at 9.2\\,$\\mu$m, which is shorter than that of many astronomical objects. It provides a good place to compare the disk properties between massive and lower-mass stars. ",
        "introduction": "In these two decades, many circumstellar disks around young forming/formed stars less than several solar masses are found by direct images in the infrared/visible observations (e.g. McCaughrean \\& O'dell 1996; Fukagawa et al. 2004; Fujiwara et al. 2006; Lagage et al. 2006; see also web database \\footnote{A comprehensive list of spatially resolved disks is available (www.circumstellardisks.org).}). Radio, infrared, and visible line observations confirmed that some disks have rotating disk kinematics as expected (Simon, Dutrey, \\& Guilloteau 2000; Pontoppidan et al. 2008; Acke, van den Ancker, Dullemond 2005). In contrast, the formation scenario for massive stars ($>$8M$_\\sun$) is still unclear. Since, for such massive stars, the time scale for the Kelvin-Helmholtz contraction is shorter than that of free-fall or accretion with accretion rate similar to low mass star cases, they start releasing energy through nuclear fusion even during accretion (Palla \\& Stahler 1993). Then radiation pressure due to their large luminosities may prevent surrounding material from accreting onto the star, in particular, in the case of very massive stars (Kahn 1974; Wolfire \\& Cassinelli 1987). Several ideas are proposed to overcome the problem: mass accretion through circumstellar disks (Yorke \\& Sonnhalter 2002; Krumholz et al. 2009), mass accretion under a much larger accretion rate than that usually considered for low mass stars (McKee \\& Tan 2002; Krumholz et al. 2009), or merging of low mass stars (Bonnell, Vine, \\& Bate 2004). Among these ideas, non-isotropic accretion through their circumstellar disks seems most plausible at present. Such non-isotropic accretion alleviates effective radiation pressure on the accreting material. Supporting evidence for this disk scenario is recent discoveries of rotating gas fragments around possible massive young stellar objects (YSOs) by interferometric observations in the radio, especially in millimeter and submillimeter wavelength regions (Cesaroni et al. 2007). Some of them have a velocity gradient in a direction orthogonal to molecular outflow lobes, which suggests that the gas fragments rotate around the central YSOs (Zhang, Hunter, Sridharan 1998; Cesaroni et al. 2005). About a dozen of objects as candidate disks around massive YSOs of up to 20\\,M$_\\sun$ are found so far (Zhang et al. 1998; Patel et al. 2005; Cesaroni et al. 2006; Beltr\\'{a}n et al. 2006; Cesaroni et al. 2007). Typically, their estimated stellar mass, luminosity, and disk radius are 4 to less than about 20\\,M$_\\odot$, a few $\\times$(10$^3$--10$^4$)\\,L$_\\odot$, and 500--2000\\,AU, respectively. It is suggested that early B Herbig Be stars are surrounded by flattened structures from measurement of depolarization across H$\\alpha$ line although the discussed scale is much smaller (order of several stellar radii) than the disk size indicated above (Vink et al. 2002).  While radio interferometric observations have so far been the most successful in unveiling the disk existence around massive YSOs, the resolution around 1$''$ is not sufficient to draw the detailed disk geometry. It is in contrast to the situation that disks around young forming stars less than several solar masses have been well depicted by direct images in visible to infrared wavelengths. For some massive YSOs, existence of disks is discussed from the polarization vector distribution of infrared scattered light image of the outflow cavities (Jiang et al. 2005), but their disks themselves are not seen because such objects are still embedded deeply in their envelopes. Direct images are strongly required to establish the existence and shape of the disks around massive YSOs. Our new approach is to search for disks in the mid-infrared around massive YSOs that have emerged from their natal clouds. Owing to the large luminosities of the central stars, the disk surface can be heated up out to large radii enough to be resolved in the mid-infrared with 8\\,m-class telescopes, which provide 100\\,AU resolution for nearby ($\\sim$400\\,pc) targets. We carry out survey observations for extended emission around Herbig Be stars and report the discovery of a disk around HD200775 in this paper. ",
        "conclusions": "We have made mid-infrared imaging and spectroscopy of the Herbig B3e star HD200775 of $\\sim$10\\,M$_\\sun$ with Subaru/COMICS. We found elongated emission around the unresolved excess emission in all of the 8.8, 11.7, 18.8, and 24.5$\\mu$m bands. The elongated elliptical emission has a radius of about 1000\\,AU and configuration perpendicular to the CO cavity wall formed by the past outflow and parallel to the projected major axis of the central close binary orbit. The brightness contours of the elliptical emission are shifted from the stellar position or from the unresolved source position. All of these observed characteristics are well explained as the surface emission of the tilted flared circumbinary disk. The present observations provide the first well-resolved infrared images of a disk around 10\\,M$_\\sun$ YSO and strongly support that HD200775 is formed through the disk accretion. The optical depth and the silicate features of the disk emission support the 'disk atmosphere' configuration, which is well modeled for low to intermediate mass stars. We fit the observed image with a flared disk model. The overall observed structure is well explained by the model. The derived flared geometry is much flatter than the disk around the intermediate mass star HD97048. The value of $z(r)/r$ is estimated to be 0.12 to 0.27 for the HD200775 disk, while it is 0.49 for the HD97048 disk. It supports the idea that there is a qualitative difference between disks around massive stars and lower mass stars. The 3.4\\,mm free-free emission image in literature is very similar to the mid-infrared elliptical disk emission in size and shape. It suggests that the disk surface photoevaporates due to the ionizing photons from the central star. The SMA 350\\,GHz observations detect the unresolved emission with a 0.8$''$ beam. The disk mass concentrated around the star is estimated as around 0.02\\,M$_\\sun$, which is similar to that of the minimum mass the solar nebula. The 10\\,$\\mu$m region spectra of the peak unresolved source and at the 1.3$''$ south are discussed. The spectrum of the unresolved source shows 1600\\,K featureless emission, which correspond to the innermost circumstellar disk in the vicinity of the central star(s). The spectrum of the diffuse emission at 1.3$''$ south is dominated by the amorphous silicate emission. The peak at 9.2\\,$\\mu$m is shorter than the usual silicate features and unique to the HD200775 disk. It may suggest alternation of grains due to plasma irradiation."
    },
    "0910/0910.1211_arXiv.txt": {
        "abstract": "{Recent spectro-polarimetric observations have provided detailed measurements of magnetic field, velocity and intensity during events of magnetic field intensification in the solar photosphere.} {By comparing with synthetic observations derived from MHD simulations, we aim to discern the physical processes underlying the observations, as well as to verify the simulations and the interpretation of the observations.} {We consider the temporal evolution of the relevant physical quantities for three cases of magnetic field intensification in a numerical simulation. In order to compare with observations, we calculate Stokes profiles and take into account the spectral and spatial resolution of the spectropolarimeter (SP) on board Hinode. We determine the evolution of the intensity, magnetic flux density and zero-crossing velocity derived from the synthetic Stokes parameters, using the same methods as applied to the Hinode/SP observations to derive magnetic field and velocity information from the spectro-polarimetric data. } {The three events considered show a similar evolution: advection of magnetic flux to a granular vertex, development of a strong downflow, evacuation of the magnetic feature, increase of the field strength and the appearance of the bright point. The magnetic features formed have diameters of 0.1-0.2\\arcsec. The downflow velocities reach maximum values of 5-10 km/s at $\\tau=1$.  In the largest feature, the downflow reaches supersonic speed in the lower photosphere. In the same case, a supersonic upflow develops approximately $200$~s after the formation of the flux concentration. We find that synthetic and real observations are qualitatively consistent and, for one of the cases considered, agree very well also quantitatively. The effect of finite resolution (spatial smearing) is most pronounced in the case of small features, for which the synthetic Hinode/SP observations miss the bright point formation and also the high-velocity downflows during the formation of the smaller magnetic features.} {The observed events are consistent with the process of field intensification by flux advection, radiative cooling, and evacuation by strong downflow found in MHD simulations. The quantitative agreement of synthetic and real observations indicates the validity of both the simulations and the interpretations of the spectro-polarimetric observations.} ",
        "introduction": "Magnetic field is ubiquitously present in the solar photosphere \\citep{deWijn:etal:2008}. On granular scales, it undergoes continual deformation and displacement. It is swept by the horizontal flows and concentrated in the intergranular lanes. Flows are able to compress the field so that the magnetic energy density $B^{2}/8\\pi$ approaches the kinetic energy density $\\rho v^{2}/2$ of the flow \\citep{Parker:1963,Weiss:1966}. This results in a magnetic field strength of a few hundred Gauss at the solar surface. Further intensification to kG strength is driven by the mechanism referred to as: \\textit{superadiabatic effect} \\citep{Parker:1978}, \\textit{convective collapse} \\citep{Webb:Roberts:1978,Spruit:Zweibel:1979} or \\textit{convective intensification} \\citep{GrossmannDoerth:etal:1998}.  The first two concepts are a theoretical idealization of the process. The superadiabatic effect contains the basic idea. \\citet{Parker:1978} pointed out that a thermally isolated dowflowing gas within the flux tube in a superadiabatically stratified environment will be accelerated, which would lead to evacuation of the flux tube. Because of the resulting pressure deficit, the gas inside the flux tube will then be pressed together (together with the frozen-in magnetic field) by the surrounding gas, causing the magnetic pressure to increase until a balance of total pressure (magnetic + gas) is reached.  The convective collapse extents the concept to the convective instability. It starts with a flux tube in thermal and mechanical equilibrium with the surrounding hydrostatically superadiabaticly stratified plasma. Since external stratification is convectively unstable, any vertical motion within the flux tube can be amplified. Downward flow will grow in amplitude and drain the material from the flux tube. The process continues until a new equilibrium with a strong field is reached. Different aspects of the concept have been the subject of extensive research \\citep[see][for reviews]{Schuessler:1990,Steiner:1999}. \\begin{figure*} \\centering \\includegraphics[width=0.95\\hsize,]{fig0.ps} \\caption{Maps of the whole simulation domain at $t=140$~s. Normalized continuum intensity at $630$~nm(left), vertical component of magnetic field (middle) and velocity (right) at a geometrical height roughly corresponding to the level of $\\langle\\tau_{500}\\rangle = 1$ ($\\approx 930$~km above the bottom of the computational box).} \\label{fig:snap} \\end{figure*} The term convective intensification is used for magnetic field intensification in realistic MHD simulations, where the process occurs in its full complexity. It is driven by the thermal effect in the surface layer of the magnetic concentration. There, due to the presence of magnetic field, heat transport by convection is reduced. The material inside the concentration radiates more that it receives. This leads to cooling of material which thus starts to sink and partial evacuation of the concentration occurs. Contraction of the magnetic concentration by the surroundings (result of the pressure imbalance) leads to an increase in magnetic field strength. Thus, the simulations \\citep{Nordlund:1983,GrossmannDoerth:etal:1998,Gadun:etal:2001,Voegler:etal:2005,Cheung:etal:2008} bear out the basic properties described by idealized concepts. That is the downflow, the evacuation of the magnetic structure, the field increase and, in some cases, establishment of a new equilibrium. The 3D MHD simulations show that the strong magnetic concentrations form as the horizontal flows in the intergranular lanes advect weak, nearly vertical field and concentrate it at the vertices of granular and mesogranular downflow lanes \\citep{Stein:Nordlund:1998,Stein:Nordlund:2006}. Larger magnetic structures form at sites where a granule submerges and the surrounding field is pushed into the resulting dark region. Whether the formed concentration appears dark or bright in the continuum intensity depends on whether the vertical cooling is compensated or not by the lateral heating due to horizontal energy exchange \\citep{Bercik:etal:2003,Voegler:etal:2005}. This formation scenario is consistent with the observations described by \\citet{Muller:1983} and \\citet{Muller:Roudier:1992}. Their observations show that network bright points form in intergranular spaces, at the junction of converging granules as the magnetic field gets compressed by the converging granular flow.    The 2D simulations by \\citet{GrossmannDoerth:etal:1998} revealed that magnetic flux concentrations formed by convective intensification can evolve in different ways. They present two possible outcomes. Depending on the initial magnetic flux, a magnetic concentration can reach a stable state after the process, or can be dispersed due to an upflow that develops as high speed downflowing material rebounces from the dense bottom of the tube. Similar results were presented by \\citet{Takeuchi:1999} and \\citet{Sheminova:Gadun:2000}.  Observational evidence was found for both cases. \\citet{BelloGonzalez:etal:2008} reported on the formation of a magnetic feature at the junction of intergranular lanes, without any significant upflow observed. \\citet{BellotRubio:etal:2001}, on the other hand, detected a strongly blueshifted Stokes V profile originating in a upward propagating shock, $13$ minutes after the amplification of magnetic field. \\citet{SocasNavarro:MansoSainz:2005} found that supersonic upflows are actually quite common.  Events that are interpreted as convective collapse were detected also with the spectropolarimeter (SP) \\citep{Lites:etal:2001} of the Solar Optical telescope \\citep{Tsuneta:etal:2008} on board Hinode \\citep{Kosugi:etal:2007}. Both, \\citet{Shimizu:etal:2008} and \\citet{Nagata:etal:2008} show cases of high speed downflows followed by magnetic field intensification and bright point appearance. The event described by \\citet{Nagata:etal:2008} shows stronger field strength and upflow in the final phase of evolution.  In this paper we give three examples of magnetic field intensification from MURaM simulations and make detailed comparison with the results of \\citet{Nagata:etal:2008} and \\citet{Shimizu:etal:2008}.  \\begin{SCfigure*} \\centering \\includegraphics[width=0.75\\textwidth]{fig1_g25l.ps} \\caption{Evolution of the continuum intensity at $630$~nm and the magnetic field in a $4\\arcsec\\times 3\\arcsec$ sized region. \\textit{Left-hand side} (double column): original spatial resolution; red and blue contours outline the downflow of 6~km/s at 80~km and upflow of 4~km/s at 400~km above $\\tau = 1$, respectively; horizontal lines mark the positions of vertical cuts shown in Figs.~\\ref{fig:cut} and ~\\ref{fig:cut_f23}; boxes corresponding to cases I, II and III at coordinates $[2\\arcsec,1.5\\arcsec]$, $[0.5\\arcsec,1\\arcsec]$ and $[3.5\\arcsec,1\\arcsec]$, respectively, enlarged in Fig.~\\ref{fig:evol_zoom}. \\textit{Right-hand side}: synthetic Hinode/SP observations; left column shows continuum intensity while the right column shows the apparent field strength (see the text); red contour marks the region with $0.01$~pm of signal excess (see text); yellow crosses mark the positions of pixels that we studied in detail in Fig.~\\ref{fig:evol_red}. Grey scales cover the range of 0.6-1.6 and 0.8-1.2 for normalized intensity at original and reduced resolution, respectively, $\\pm 1000$~G for the vertical component of magnetic field (original resolution) and $\\pm200$~Mx/cm$^{2}$  for the apparent longitudinal magnetic flux density (Hinode resolution).} \\label{fig:evol} \\end{SCfigure*} ",
        "conclusions": "Our case study of three examples of magnetic field intensification has shown that in all three cases, the field is advected to the junction of several granules. There, it is confined by converging granular flows, which, in two cases, form a vortex. Owing to the presence of the magnetic field, the thermal effect (radiative cooling) induces evacuation of the flux concentration. The evacuation leads to a downward shift of the optical depth scale within the flux concentrations. The shift is smaller for the smaller features due to the lateral radiative heating, which inhibits further evacuation \\citep{Venkatakrishnan:1986}. As a result, the magnetic field at $\\tau=1$, in the smaller features, is weaker than in the case of the feature with more flux. This is in accordance with recent numerical \\citep{Cheung:etal:2007} and observational \\citep{Rueedi:etal:1992,Solanki:etal:1996} results.  During the evacuation, the dowflow velocities reach maximum values of 5-10 km/s at $\\tau=1$. In the case of the largest feature, the downflow extends from the upper boundary of the simulation domain and becomes supersonic in the lower photosphere. The magnetic features formed have diameters of 0.1-0.2~\\arcsec. In the case of the biggest feature, a supersonic upflow develops approximately $200$~s after the formation of the flux concentration. The upflow does not lead to a complete dispersal of the field, but the feature persists until 9 minutes later when it undergoes a fragmentation. The disappearance of the features in all three cases occurs when they meet opposite polarity features, between 3 and 20 minutes after their formation.  We also show what happens with the observables when the effects of smearing to observational spatial resolution is taken into account. An important result is that, in the case of small features, Hinode/SP would miss the bright point formation and, in some cases, also the high velocity downflows that develop in the process. On the other hand, the signatures of the evolution of large features are detectable even after the spatial smearing. We show that this case can be quantitatively compared with Hinode/SP observation \\citep{Nagata:etal:2008,Shimizu:etal:2008} and exhibits a very similar evolution. This suggests that the magnetic field intensification process in the MURaM simulations is a faithful description of the process taking place on the Sun. Furthermore, our study indicates that the analysis and interpretation of the observations in terms of the convective intensification process is well-founded."
    },
    "0910/0910.1357_arXiv.txt": {
        "abstract": "{I present a predictive analysis for the behavior of the far-infrared (FIR)--radio correlation as a function of redshift in light of the deep radio continuum surveys which may become possible using the square kilometer array (SKA). To keep a fixed ratio between the FIR and predominantly non-thermal radio continuum emission of a normal star-forming galaxy, whose cosmic-ray (CR) electrons typically lose most of their energy to synchrotron radiation and Inverse Compton (IC) scattering, requires a nearly constant ratio between galaxy magnetic field and radiation field energy densities. While the additional term of IC losses off of the cosmic microwave background (CMB) is negligible in the local Universe, the rapid increase in the strength of the CMB energy density (i.e. $\\sim(1+z)^{4})$ suggests that evolution in the FIR-radio correlation should occur with infrared (IR;~$8-1000~\\mu$m)/radio ratios increasing with redshift. This signature should be especially apparent once beyond $z\\sim3$ where the magnetic field of a normal star-forming galaxy must be  $\\sim$50~$\\mu$G to save the FIR-radio correlation. At present, observations do not show such a trend with redshift; $z\\sim6$ radio-quiet quasars appear to lie on the local FIR-radio correlation while a sample of $z\\sim4.4$ and $z\\sim2.2$ submillimeter galaxies (SMGs) exhibit ratios that are a factor of $\\sim$2.5 {\\it below} the canonical value. I also derive a 5$\\sigma$ point-source sensitivity goal of $\\approx$20~nJy (i.e. $\\sigma_{\\rm RMS} \\sim 4$~nJy) requiring that the SKA specified be $A_{\\rm eff}/T_{\\rm sys}\\approx  15000$~m$^{2}$~K$^{-1}$; achieving this sensitivity should enable the detection of galaxies forming stars at a rate of $\\gtrsim25~M_{\\odot}~{\\rm yr}^{-1}$, such as typical luminous infrared galaxies (i.e. $L_{\\rm IR} \\gtrsim 10^{11}~L_{\\odot}$), at all redshifts if present. By taking advantage of the fact that the non-thermal component of a galaxy's radio continuum emission will be quickly suppressed by IC losses off of the CMB, leaving only the thermal (free-free) component, I argue that deep radio continuum surveys at frequencies $\\gtrsim$10~GHz may prove to be the best probe for characterizing the high-$z$ star formation history of the Universe unbiased by dust.  }  \\FullConference{Panoramic Radio Astronomy: Wide-field 1-2 GHz research on galaxy evolution\\\\ June 2-5 2009\\\\ Groningen, the Netherlands}  \\begin{document} ",
        "introduction": "Radio continuum emission from galaxies arises due to a combination of thermal and non-thermal processes primarily associated with the birth and death of young massive stars, respectively. The thermal (free-free) radiation of a star-forming galaxy is emitted from H{\\sc ii} regions and is directly proportional to the photoionization rate of young massive stars. Since emission at GHz frequencies is optically thin, the thermal radio continuum emission from galaxies is a very good diagnostic of a galaxy's massive star formation rate. Massive ($\\gtrsim 8~M_{\\odot}$) stars which dominate the Lyman continuum luminosity also end their lives as supernovae (SNe) whose remnants (SNRs) are responsible for the acceleration of cosmic-ray (CR) electrons into a galaxy's general magnetic field resulting in diffuse synchrotron emission. These same massive stars are often the primary sources of dust heating in the interstellar medium (ISM) as their starlight is absorbed and reradiated at far-infrared (FIR) wavelengths by interstellar grains. This common origin between the FIR dust emission and thermal $+$ non-thermal radio continuum emission from galaxies is thought to be the dominant physical processes driving the FIR-radio correlation on global (e.g. \\cite{gxh85,yrc01}, and references therein) and local (e.g. \\cite{ejm06,ejm08}, and reference therein) scales. At present,  all indications suggest that the FIR-radio correlation holds out to moderate redshifts (e.g. \\cite{df06,ejm09a,ms09}, and references therein).   The detection of large populations of dusty star-forming galaxies at high redshift by ISO and {\\it Spitzer} has underscored the need for reliable star formation rate diagnostics unaffected by dust. While radio emission may provide an excellent advantage over other wavelengths, detecting large populations of high redshift star-forming galaxies at radio wavelengths remains extremely difficult. While detectable with current FIR capabilities, even with a fully operational EVLA, IR-bright star-forming galaxies  (e.g. M~82; $L_{\\rm IR} \\approx 4\\times10^{10}~L_{\\odot}$) and moderate LIRGs  (i.e. $L_{\\rm IR} \\approx 3\\times10^{11}~L_{\\odot}$) will not be detectable beyond redshifts of $z\\sim 1$ and $z\\sim 2$, respectively. A next-generation radio facility such as the Square Kilometer Array (SKA) should easily remedy this disparity between the depth of FIR and radio continuum surveys.  Recently, \\cite{ejmc09} described physically motivated predictions for the evolution of the FIR-radio correlation as a function of redshift arising from variations in the CR electron cooling time-scales as Inverse Compton (IC) scattering off of the Cosmic Microwave Background (CMB) becomes increasingly important. Since the non-thermal component of a galaxy's radio continuum is increasingly suppressed with redshift, radio continuum measurements at moderately high frequency ($\\sim$10~GHz) become one of the cleanest ways to quantify the star formation activity of galaxies at high redshifts unbiased by dust. In this proceedings article I summarize some of these findings with an emphasis placed on what this might mean for deep radio continuum surveys using the SKA. ",
        "conclusions": "In this proceedings article I have summarized the conclusions of \\cite{ejmc09} which presented a predictive analysis for the expected evolution of the FIR-radio correlation versus redshift arising from variations in the CR electron cooling time-scales as IC scattering off of the CMB becomes increasingly important. In doing so, I have focussed on the value of deep radio continuum surveys in studies of star formation at high-$z$, particularly in the context of the SKA, finding the following:  \\begin{enumerate}  \\item Deep radio continuum observations at frequencies $\\gtrsim$10~GHz using next generation facilities like the SKA will likely provide the most accurate measurements for the star formation rates of normal galaxies at high~$z$. The non-thermal emission from such galaxies should be completely suppressed due to the increased IC scattering off of the CMB leaving only the thermal (free-free) emission detectable; this situation may be complicated if `anomalous' microwave emission from spinning dust grains \\cite{dl98a} in such systems is not negligible.  \\item For normal star-forming galaxies to remain on the local FIR-radio correlation at high redshifts requires extraordinarily large magnetic field strengths to counter IC losses from the CMB. For example, the magnetic field of a $z\\sim6$ galaxy must be $\\gtrsim$500~$\\mu$G to obtain a nominal IR/radio ratio. Thus, galaxies which continue to lie on the FIR-radio correlation at high-$z$, such as the sample of radio-quiet QSO's, most likely have radio output which is dominated by an AGN, or the physics at work is simply unknown.   \\item To detect typical LIRGs ($L_{\\rm IR} \\gtrsim 10^{11}~L_{\\odot}$) at all redshifts will require nJy sensitivities at GHz frequencies, specifically a 5$\\sigma$ point-source sensitivity of $\\approx$20~nJy (i.e. $\\sigma_{\\rm RMS} \\approx 4$~nJy). Thus, for the SKA to achieve this sensitivity for a reasonably long (300~hr) integration necessitates that $A_{\\rm eff}/T_{\\rm sys} \\sim 15000$~m$^{2}$~K$^{-1}$. At this sensitivity, the SKA will be sensitive to all galaxies forming stars at $\\sim$25~$M_{\\odot}~{\\rm yr}^{-1}$. This includes a significant amount of sources included in the UV luminosity function work of \\cite{rb07,rb08} between $4\\lesssim z \\lesssim 9$.  \\end{enumerate}"
    },
    "0910/0910.0398_arXiv.txt": {
        "abstract": "{ Supersonic turbulence in molecular clouds is a dominant agent that strongly affects the clouds' evolution and star formation activity. Turbulence may be initiated and maintained by a number of processes, acting at a wide range of physical scales. By examining the dynamical state of molecular clouds, it is possible to assess the primary candidates for how the turbulent energy is injected. } { The aim of this paper is to constrain the scales at which turbulence is driven in the molecular interstellar medium, by comparing simulated molecular spectral line observations of numerical magnetohydrodynamic (MHD) models and molecular spectral line observations of real molecular clouds.} { We use principal component analysis, applied to both models and observational data, to extract a quantitative measure of the driving scale of turbulence. } { We find that only models driven at large scales (comparable to, or exceeding, the size of the cloud) are consistent with observations. This result applies also to clouds with little or no internal star formation activity.} { Astrophysical processes acting on large scales, including supernova-driven turbulence, magnetorotational instability, or spiral shock forcing, are viable candidates for the generation and maintenance of molecular cloud turbulence. Small scale driving by sources internal to molecular clouds, such as outflows, can be important on small scales, but cannot replicate the observed large-scale velocity fluctuations in the molecular interstellar medium.} ",
        "introduction": "Turbulence is an important agent that controls the evolution (and perhaps formation) of molecular clouds and the subsequent production of stars. As such, it has attracted significant attention from theorists, especially since the advent of numerical supercomputer simulations. Of particular interest is the source(s) of energy injection that create and sustain turbulence in molecular clouds. A number of different mechanisms have been proposed, including supernovae, H{\\sc II} regions, outflows, spiral arms, magneto-rotational instability in galactic disks (Mac Low \\& Klessen 2004; Miesch \\& Bally 1994). These mechanisms may be distinguished by the effective spatial scale at which they preferentially operate, and clues to the nature of the energy injection mechanism(s) may be extracted from spectral line imaging observations of molecular clouds.  A number of methods for studying resolved velocity fields in molecular clouds have been developed and applied. These include projected velocity (line centroid) analysis (e.g. Scalo 1984; Miesch \\& Bally 1994; Ossenkopf \\& Mac Low 2002; Brunt \\& Mac Low 2004), the spectral correlation function (SCF; Rosolowsky {\\it et al}{} 1999), velocity channel analysis (VCA; Lazarian \\& Pogosyan 2001, 2004), and principal component analysis (PCA; Heyer \\& Schloerb 1997). To date, these methods have been used to estimate the power law indices of the velocity structure function/power spectrum in molecular clouds from observed data cubes of molecular line emission  (e.g. Brunt \\& Heyer 2002b; Heyer \\& Brunt 2004).  Application of PCA to Outer Galaxy molecular clouds (Brunt 2003a -- Paper I hereafter) revealed that, in comparison to simple models, the observational record favoured large scale driving of turbulence in the molecular clouds. In their study of the Polaris molecular cloud, Ossenkopf \\& Mac Low (2002) also found that large scale driving of turbulence provided a better explanation of the cloud's velocity structure.  In this paper, we construct simulated observations of molecular clouds, derived from computational simulations of interstellar turbulence. The models include magnetic fields and self-gravity and are driven (randomly forced) on a range of spatial scales. We employ PCA to quantitatively investigate the observational signatures of different driving scales. Our numerical measurements are compared to previous PCA results obtained from the simple cloud models of Paper I and to the same measurements made on real molecular clouds. The layout of the paper is as follows. In Section~2, we briefly summarize the PCA method and review the relevant findings of Paper I. Section~3 introduces the numerical models and summarizes the simulated observations of these. In Section~4, we present our results, compare these to corresponding observations, and discuss the implications for the generation of turbulence in molecular clouds. Our conclusions are given in Section~5. ",
        "conclusions": "}}  Both analytical and computational descriptions of turbulence are necessarily constrained by observations of interstellar clouds. A qualitative inspection of Figure~\\ref{fig:pcseq} shows that the eigenimages derived from clouds models with large $\\lambda_{D}$/L$_{c}$ are more consistent with the observations of NGC~7538. More generally, the measured values of $l_{2}$/$l_{1}$ from real molecular clouds are typically $\\gtrsim$~0.2. In Figure~\\ref{fig:lhist} we plot the histogram of $l_{2}$/$l_{1}$ measured in the sample of clouds from Paper I, to which we have added additional measurements from the clouds analyzed in Heyer \\& Brunt (2004). In the combined sample there are 35 clouds in total. Using Figure~\\ref{fig:lolc} as a guide to the relationship between $<l_{2}$/$l_{1}>$ and $\\lambda_{D}$/L$_{c}$, these values imply that the molecular clouds are dominated by turbulence driven on large scales compared to the cloud sizes. This may be simply a result of the driving scale itself determining the size of molecular clouds (Ballesteros-Paredes \\& Mac Low 2002; Paper I).  \\begin{figure} \\centering \\includegraphics[width=9cm]{f3.eps} \\caption{A comparison of velocity power spectra (power $P$ versus wavenmuber $k$) obtained from a numerically simulated cloud (HC8, with $k_{d}$~=~7--8) and an fBm field ($k_{cut}$~=~7) of the type used in Paper I to represent turbulent driving at $k$~$\\approx$~7. HC8 has more power at low wavenumbers relative to the fBm field. (The vertical scale in this plot is arbitrary.) } \\label{fig:powerspec} \\end{figure}  \\begin{figure} \\centering \\includegraphics[width=9cm]{f4.eps} \\caption{(a) Comparison of $l_{2}$/$l_{1}$ derived from simulated CO observations and observations using the $v$-hist method where opacity and excitation effects are not included. (b) Comparison of $l_{2}$/$l_{1}$ derived from simulated $^{13}$CO and $^{12}$CO observations. For each p lot, the solid line denotes equivalent values along the ordinate and absissca axes. } \\label{fig:isotope} \\end{figure}  In our experiment, we have considered the simplified case where a single ``driving scale'' is in operation. Within this limitation we identify large scale driving as the dominant scenario. In reality, turbulence can in principle be driven on multiple scales by a number of mechanisms (Scalo 1987). The origin of large-scale energy injection is discussed by Mac Low \\& Klessen (2004), who concluded that field supernovae were the dominant mechanism in regions where they occur, while magneto-rotational instability (Kim, Ostriker, \\& Stone 2003; Tamburro {\\it et al}{} 2009) may provide a background level. In addition to these, other possible mechanisms include forcing by shocks in spiral arm potentials; Dobbs \\& Bonnell (2007) demonstrate that the scale-dependent velocity dispersion in molecular clouds can be replicated by simulated clouds in a galactic disk with a fixed spiral arm pattern. Most of these processes likely require that the molecular cloud turbulence is inherited from still larger scale motions in the atomic ISM (Elmegreen 1993, Ballesteros-Paredes {\\it et al} 1999; Brunt 2003a). In this scenario, the ``driving'' of molecular cloud turbulence could simply be due to the continuous downward cascade of turbulent energy, that not only injects the turbulence but is also responsible for the (potentially rapid) molecular cloud formation in the first place (Bergin {\\it et al} 2004; Glover \\& Mac Low 2007). The presence of large scale turbulence in molecular clouds would be a natural, inevitable consequence of their formation, and their subsequent evolution can be significantly affected by dynamical events occurring in the larger scale ISM.  \\begin{figure} \\centering \\includegraphics[width=9cm]{f5.eps} \\caption{Histogram of $l_{2}$/$l_{1}$ obtained from $^{12}$CO observations of real molecular clouds. The vertical lines mark the mean $l_{2}$/$l_{1}$ derived from the model observations ($^{12}$CO) and the horizontal arrows extend over the range of measured $l_{2}$/$l_{1}$. } \\label{fig:lhist} \\end{figure}  Energy injection on (initially) small scales by the spatio-temporally intermittent development of outflows, stellar winds and H{\\sc ii} regions within the cloud may not be well modeled by random forcing methods used in these simulations. These point-like injections of energy can expand their spheres of influence over time and may ultimately contribute to large scale turbulent motions. However, on large scales, these processes are disfavoured on energetic grounds (Mac Low \\& Klessen 2004). While there is evidence that energy injection by outflows can be important over limited scales (e.g. Bally, Devine, \\& Alten 1996; Knee \\& Sandell 2000) it is unlikely that outflow-driven turbulence can explain the origin of molecular cloud turbulence as a whole (Walawender, Bally, \\& Reipurth 2005; Banerjee, Klessen, \\& Fendt 2007). This is demonstrated by recent simulations of outflow-driven turbulence which reveal that energy injection by outflows is not capable of creating turbulence at scales comparable to the cloud size. Models of interacting outflows generated either randomly (Carroll {\\it et al}{} 2008), or self-consistently (Nakamura \\& Li 2007) show that turbulence is only injected with an effective driving scale of about 1/5 to 1/10 the size of the cloud ($\\lambda_{D}/L_{c}$~$\\approx$~0.1--0.2) which is incompatible with our results as summarized in Figure~\\ref{fig:lolc} and Figure~\\ref{fig:lhist}. The observable ratio $l_{2}$/$l_{1}$ is expected to lie in the range 0.02--0.05 when $\\lambda_{D}/L_{c}$~$\\approx$~0.1--0.2, according to our modelling results. Additionally, the cloud modelled by Nakamura \\& Li (2007) is only 1.5~pc in size, and it is unclear whether the effective driving scale would increase (for the same outflow parameterization) if a larger cloud was modelled. If the {\\it fractional} driving scale of 0.1--0.2 is interpreted as a {\\it physical} driving scale of 0.15--0.3~pc, then outflow-driven turbulence would be even less effective in globally exciting turbulence in larger clouds. On the other hand, in larger clouds, more massive and energetic outflows may be expected to be present, but it is not currently clear how (or if) the effective fractional driving scale would increase.  An observational estimate of the effective driving scale of turbulence by outflows was found by Swift \\& Welch (2008). They inferred an energy injection scale of 0.05~pc for L1551, which is a small fraction of the the overall cloud diameter of around 1.8~pc. Using this estimate, they found a rough balance between the energy injection rate (from the outflows) and the turbulent dissipation rate, with a characteristic injection/decay timescale of $\\sim$~0.1~Myr, which is substantially less than the inferred cloud age of $\\sim$~4--6~Myr. We note here that some caution is required in interpreting the appearance of injection/decay balance for the outflow-driven turbulence. The dissipation time scale of turbulence is proportional to the driving scale (Mac Low 1999). Swift \\& Welch (2008), in calculating their dissipation rate, used a driving scale of 0.05~pc, and therefore their result shows primarily that the energy injection through outflows is quickly dissipated on short time scales over short length scales. If the L1551 cloud is, or has been, subject to large scale ($\\gtrsim$~1.8~pc) driving of turbulence, then the large scale turbulence is controlled by a much longer dissipation timescale of (1.8/0.05)~$\\times$~0.1~Myr~$\\approx$~3.6~Myr, which is more in line with the cloud age. The small-scale driving of outflows would then occur within the longer time evolution of the cloud, set by the longer dissipation time scales of the initial turbulence, injected on large scales.   \\begin{figure} \\centering \\includegraphics[width=9cm]{f6.eps} \\caption{ First and second eigenmages obtained from principal component analysis of the NGC~1333 molecular cloud. The C$^{18}$O data and analysis are confined to the central core region, delineated by the rectangular box on the $^{12}$CO and $^{13}$CO images.} \\label{fig:eigim1333} \\end{figure}  The effect of multiple outflows within a small region of space may be seen in the outflow-rich NGC~1333 molecular cloud. Here, Quillen {\\it et al} (2005) describe as many 22 cavities within a 1~pc$^{3}$ volume, possibly excavated by outflow activity, in the $^{13}$CO (J=1--0) map of Ridge {\\it et al} (2003). The cavities have typical diameters of 0.1--0.2~pc, indicative again of a small effective driving scale, as the shells surrounding the cavities would presumably collide and merge at larger scales. However, it is not yet clear that outflows are the driving source for these cavities, as many do not have obvious stellar sources inside them -- see Quillen {\\it et al} (2005) for further discussion.  To investigate outflow-driven turbulence from an observational perspective, we applied the PCA method to CO observations of the NGC~1333 molecular cloud. We used the J=1--0 spectral lines of $^{12}$CO and $^{13}$CO observed at FCRAO as part of the COMPLETE project (Ridge {\\it et al} 2006), as well as FCRAO C$^{18}$O J=1--0 spectral line data towards the central core region of NGC~1333. In Figure~\\ref{fig:eigim1333} we show the first two eigenimages obtained from the analysis for each spectral line. For $^{12}$CO and $^{13}$CO, we find ``dipole'' second eigenimage structure characteristic of large scale turbulence, and measure $l_{2}$/$l_{1}$ values of 0.59 ($^{12}$CO) and 0.63 ($^{13}$CO). The overall cloud size is estimated from the $^{12}$CO $l_{1}$ measurement to be 3.27~pc, assuming a distance of 318~pc. These measurements show that turbulence is (or has been) driven on large scales in NGC~1333, and is unlikely to have originated from the outflows, which are confined to the central core region, marked by the small box in Figure~\\ref{fig:eigim1333}. Analysis of the C$^{18}$O data in this box allows us to focus in on the high column density material lying in the immediate vicinity of the outflows. We measure $l_{2}$/$l_{1}$~=~0.18$\\pm$~0.07 for the high column density material traced by C$^{18}$O, which is substantially smaller than the global $l_{2}$/$l_{1}$ values found using $^{12}$CO and $^{13}$CO, but still reasonably consistent with turbulence driven at large scales. Some caution should be applied to this result, because, as noted above, large temporal variations in $l_{2}$/$l_{1}$ can occur in the case of large scale driving. With this proviso, according to our model results, the measured $l_{2}$/$l_{1}$ for the central region would imply a fractional driving scale of $\\lambda_{D}/L_{c}$~$\\approx$~0.5--1.0, or a physical driving scale of 0.43--0.86~pc, based on the measured $l_{1}$~=~0.86~pc for the C$^{18}$O data.  For reference, the cavity sizes of 0.1--0.2~pc in NGC~1333, if taken as a measure of the driving scale within the 0.86~pc C$^{18}$O central core region, best match our models driven at $k_{d}$~=~3--4, for which we find $l_{2}$/$l_{1}$~$\\approx$~0.11. Examination of the C$^{18}$O second eigenimage structure reveals that it shares, to some degree, the same north-south ``dipole'' structure seen in the $^{12}$CO and $^{13}$CO second eigenimages. The presence of this signature, along with the $l_{2}$/$l_{1}$~=~0.18 measurement, suggests that both the large-scale turbulence in the cloud as a whole, and small-scale (outflow) driven turbulence are important in this region. The inferred driving scale is therefore likely an intermediate value between that arising from the outflows and that deriving from the large-scale turbulent gradient across the core region. As a caveat, we note that the $^{12}$CO and $^{13}$CO lines are likely to better trace lower density, more spatially extended material than that traced by the C$^{18}$O line, so the relationship between the gradients seen in Figure~\\ref{fig:eigim1333} may not be as obvious as we assume. If the C$^{18}$O gradient is itself caused by outflow activity, then this may indicate an interesting connection between the large and small-scale energy injection mechanisms.  \\begin{figure*} \\centering \\includegraphics[width=17cm]{f7.eps} \\caption{ Plots of $\\delta{v}$ versus $l$ using all space and velocity scales obtained from principal component analysis of $^{12}$CO, $^{13}$CO, and C$^{18}$O J=1--0 data in the NGC~1333 moleclar cloud. The solid lines in each plot show the bisector fit to all points; the fitted relationships for $^{12}$CO and $^{13}$CO are repeated as dotted lines in all plots.  } \\label{fig:dvl1333} \\end{figure*}  To examine the overall scale-dependence of turbulent motions in NGC~1333, in Figure~\\ref{fig:dvl1333} we show plots of $\\delta{v}$ versus $l$ for each isotope from their respective maps (see Brunt \\& Heyer 2000b). It is noteworthy that the $^{12}$CO and $^{13}$CO measurements conform to the typical $\\delta{v}$--$l$ relationship found by Heyer \\& Brunt (2004). Of more interest here is the increase in $\\delta{v}$ seen on scales of $\\sim$~0.1~pc in the C$^{18}$O data, relative to the overall level set by the $^{12}$CO and $^{13}$CO data for the cloud as a whole. This excess kinetic energy likely derives from the effect of outflows in the central core region of the cloud. Although the number of retrieved $\\delta{v}$--$l$ pairs is small, the data are in broad agreement with Quillen et al (2005) who estimate outflow-driven cavity sizes and velocity perturbations of $\\sim$~0.1--0.2~pc and $\\sim$~1~kms$^{-1}$ respectively. Thus internal driving of turbulence can be important in sub-parsec regions of larger clouds, where a large number of outflows can develop. The PCA results for the NGC~1333 cloud as a whole set this in context, revealing the presence of larger-scale turbulence that will evolve on longer time scales than that present in the central core region. Turbulent dissipation in the dense part can therefore be replenished not only by local sources, but by external ``driving'' by larger scale flows originating in the surrounding cloud, as part of the overall hierarchy of turbulent motions. One cannot then consider the central star-forming regions as closed systems, evolving independently of their larger scale surroundings. Our cloud sample as a whole does not support a picture in which large scale turbulent motions have decayed sufficiently so that small scale driving alone is dominant. We conclude that either the clouds are continually driven on large scales, or that most clouds are sufficiently young that the initial seeding of turbulence by the large scale flows that created the cloud has not yet dissipated. Other observations (Ossenkopf \\& Mac Low 2002; Brunt \\& Mac Low 2004) support this conclusion. It is not yet clear whether clouds are continually driven, or whether the turbulence is in a decaying state. Offner, Klein, \\& McKee (2008) find that while their simulated clouds do not readily distinguish between decaying or driven conditions, there is a marginal preference for continual driving. If clouds are driven at large scales, the turbulent dissipation time is comparable to their dynamical time (Mac Low 1999).  As noted above, the dipole pattern in the second eigenimage that is observed in all molecular clouds provides an important constraint to candidate driving sources. The dipole reflects the spatial distribution of the largest velocity differences within a cloud. One can not directly discriminate whether these velocity differences are due to shear, compressive, or expanding motions. Large scale driving can readily account for such a pattern as it directly deposits the energy at these scales. Turbulence driven on small scales can in principle provide support on larger scales (Klessen {\\it et al} 2000) and it may be possible for excess small scale energy input to drive large scale expasion motion. More generally, a cluster-forming clump could experience expansion, collapse, or perhaps oscillation about an equilibrium state, depending on how active the star formation is. However, the dipole pattern suggests a more directed flow of material, which would require the combined action of outflows to act in a preferred direction. There is a possible mechanism for outflows to be oriented in a particular direction: a strong magnetic field could result in core collapse along field lines, leading to co-oriented protostellar disks and therefore co-oriented outflows, for which some evidence is presented in Anathpindika \\& Whitworth (2008). The magnetic field strength needed to impose such directivity is likely to inhibit cluster formation, and instead promote star formation in a more distributed, quiescent mode (Heitsch, Mac Low, \\& Klessen 2001; Price \\& Bate 2008). Outflows from newborn stars and H{\\sc ii} regions can also redistribute energy from small to larger scales by driving expanding shells.  Such flows may also perturb the magnetic field that threads the molecular cloud to excite Alfv\\'en waves that can further redistribute the outflow energy. However, such activity would again require implausible coherence of location and alignment of outflows to reproduce the observed dipole pattern.  Another candidate for driving large scale turbulence ``internally'' is energy injection by H{\\sc ii} regions, as argued by Matzner (2002). However, large scale driving is applicable to molecular clouds where H{\\sc ii} regions are absent, such as G216-2.5 (Maddelena's Cloud; Heyer, Williams, \\& Brunt 2006). So while these mechanisms are no doubt present in some molecular clouds, they cannot explain molecular cloud turbulence in general and their effects will be limited to small scales. If H{\\sc ii} regions become large enough to drive large scale motions, then it is likely that the cloud will be destroyed through photoionization rather than ``driven'' (Matzner 2002; Dale {\\it et al}{} 2005).  Turbulence driven at large scales promotes star formation that is clustered, rapid, and efficient, while small scale driving tends to form stars singly, slowly, and inefficiently (Klessen {\\it et al}{} 2000). If the star formation rate can be retarded by (additional) small scale energy injection, it must do this in an environment which can be significantly (perhaps dominantly) influenced by large scale turbulent flows of material. While the large scale versus small scale driving picture can be modified by the effects of magnetic fields (Nakamura \\& Li 2008; Price \\& Bate 2008), it is in a much more dynamic way than that described by the quasistatic model (Shu, Adams, \\& Lizano 1987). For example, recent high spatial dynamic range imaging of the Taurus molecular cloud (Goldsmith {\\it et al}{} 2008; Heyer {\\it et al}{} 2008) reveal large scale, magnetically-regulated, turbulent flows of material.  In addition to energy injection, another important consideration is the dissipation of turbulence. Basu \\& Murali (2001) argue that it is difficult to reconcile the inferred heating rate arising from dissipation of turbulence with observed cloud luminosities unless the driving occurs at large scales. More recently, Pan \\& Padoan (2008) show that (assuming large scale driving) heating by turbulent dissipation can exceed cosmic ray heating, and typical temperatures of $\\sim$8.5~K can be sustained by turbulent heating alone. Since the turbulent heating rate scales as $(\\lambda_{D}/L_{c})^{-1}$,  widespread small-scale driving could lead to high cloud temperatures that are incompatible with observations for molecular clouds as a whole (although not for small sub-regions within the clouds).  We mention a note of caution regarding the results presented here. The numerical simulations of turbulence relied on random forcing (in Fourier space) to generate the turbulent driving, which does not in detail adequately represent many physical sources of energy injection. In the case of outflow-generated turbulence, considered here to be ``small scale'', it was indeed found that the turbulence was effectively driven on small scales. The close correspondence between the numerical models and the simple models of Paper I suggest also that it is not necessarily the details of the flows that are essential, but simply the range of scales on which the turbulence is present. In this sense, the modeling completed so far (Paper I and this work) adequately represent, statistically, turbulence with an outer scale that is detectable in observations. It is to be expected that more realistic driving mechanisms (e.g. as implemented by Nakamura \\& Li 2007) can be investigated in future. Finally, our results recommend that simulations of randomly forced turbulence must necessarily include large scale driving in order to replicate real molecular clouds. How this translates in detail to more realisitic driving mechanisms must be addressed in future work."
    },
    "0910/0910.5268_arXiv.txt": {
        "abstract": "The development of cosmic ray air showers can be influenced by atmospheric electric fields. Under fair weather conditions these fields are small, but the strong fields inside thunderstorms can have a significant effect on the electromagnetic component of a shower. Understanding this effect is particularly important for radio detection of air showers, since the radio emission is produced by the shower electrons and positrons. We perform Monte Carlo simulations to calculate the effects of different electric field configurations on the shower development. We find that the electric field becomes important for values of the order of 1 kV/cm. Not only can the energy distribution of electrons and positrons change significantly for such field strengths, it is also possible that runaway electron breakdown occurs at high altitudes, which is an important effect in lightning initiation. ",
        "introduction": "\\label{sec:intro}  The effect of atmospheric electric fields on the development of extensive air showers from high energy cosmic rays has not received much attention in the past. Because of the large energies of shower particles, the electric fields present in the atmosphere are generally much too small to alter the particle energies significantly. The largest fields are of the order of 1 kV/cm and only occur in thunderstorms \\cite{book}. In such fields the hadronic and muonic part of the shower are hardly affected, although a muon deficit due to increased decay rate has been reported by Alexeenko et al.~\\cite{A02}. The effects on the electromagnetic shower are much larger, but they are, as we will show in this work, local in the sense that the amount and energy distribution of electromagnetic particles quickly adapts to the local background field.  Below thunderstorms the electric field decreases, so particle detector arrays will in general not be strongly sensitive to the influence of electric fields. Mountain top experiments, however, can be very close to, or even inside thunderstorms. An increase in the air shower detection rate and the single particle count has been reported by high altitude experiment such as EAS-TOP \\cite{eastop} and the Carpet air shower array \\cite{A02}.  Effects of atmospheric electric fields are generally not included in air shower simulation codes like CORSIKA \\cite{CORSIKA}. There are, however, two topics for which we demonstrate in this paper that it is important to take electric fields into consideration.  The first is the consequence of different shower developments in strong electric field regions on the radio emission of the shower. With experiments such as LOPES \\cite{F05} and Codalema \\cite{codalema}, the technique of radio detection of air showers has become mature in the last few years. Radio antennas, with their low costs and high duty cycles, have been demonstrated to be an attractive addition to large air shower arrays like the Pierre Auger Observatory \\cite{PAO}. Already in the 1970s, experimental results by Madolesi et al.~\\cite{M74} showed excessively large radio pulses during thunderstorms. Together with a large range in measured radio intensity and the inability to filter out radio interference, the unknown effect of electrical conditions in the atmosphere was a reason to abandon the radio experiments \\cite{B77} until the development of digital radio arrays like LOFAR \\cite{lofar} revived the interest. A study of the effect of weather conditions on the radio pulse strength with LOPES data confirmed the amplification of radio pulses during thunderstorm conditions, but also showed that no effect is observed under other weather conditions \\cite{B07}. In another study \\cite{N08} it was shown that during thunderstorms the air shower arrival direction reconstructed with the LOPES radio antennas can be a few degrees off with respect to the direction reconstructed by the KASCADE particle detector array.  It is now known that the radio emission of air showers can be described in terms of coherent synchrotron radiation of shower electrons and positrons that follow curved trajectories in the Earth's magnetic field as first described in Falcke et al.~\\cite{FG03} and in more detail in Huege et al.~\\cite{HF03}. Alternatively, a macroscopic picture can be constructed in which the shower charges, that are separated in the magnetic field, support a transverse current producing radiation as proposed by Kahn \\& Lerche \\cite{KL66} and more recently by Scholten et al.~\\cite{Sch}. Although there are subtle differences between these descriptions, it is clear that an alteration in the electron and positron distribution of the shower at some altitude, can influence the power of the radio pulse in both models. In this paper we limit ourselves to a study of electric field effects on these distributions, not on the radio pulse itself, which will be the subject of a forthcoming paper.  The second issue related to the electric field is the suggestion that air showers of sufficient energy can start an avalanche of runaway electrons in thunderstorm electric fields. Ionization electrons that are produced in collisions of shower particles with air molecules are accelerated in the thunderstorm electric field and can, under the right conditions, gain enough energy to ionize further molecules, an effect described by Gurevich et al.~\\cite{G92}. In thunderstorm research the field strength that can support such avalanches is known as the threshold field, described in Marshall et al.~\\cite{MMR95}. In their work, the authors present thunderstorm measurements which show that lightning often occurs when the thunderstorm field exceeds the breakeven field, suggesting that runaway electron breakdown plays a role in lightning initiation. By providing seed electrons for avalanches, air showers from cosmic rays may play an important role in thunderstorm dynamics. Simulations by Dwyer \\cite{Dw03} have shown that the threshold field strength for an avalanche to develop is slightly higher than the breakeven field, taking into account the effects of elastic scattering.  Inside thunderstorms there have been measurements of X-ray bursts \\cite{MP85} and gamma bursts, both from Earth \\cite{G04} and space \\cite{F94}. This emission can be explained in terms of bremsstrahlung emitted by the runaway breakdown electrons \\cite{GM99, Dw08a}.  In this work, we simulate the effect of electric fields on the development of air showers with CORSIKA. In Sec.~\\ref{sec:sim} we describe the setup of our simulation and the modification that has been made for CORSIKA. In Sec.~\\ref{sec:effects} we derive some simple analytic estimates and limits to compare with the results. Simulation results are presented in Sec.~\\ref{sec:res} and we conclude with a discussion of the simulation limitations and the consequences for realistic field configurations. ",
        "conclusions": "We have conducted CORSIKA simulations for vertical and inclined showers of several energies inside electric fields with varying strengths. The evolution of the electromagnetic part of the shower does not change significantly below field strengths of 100 V/cm. For fields of the order of a 1000 V/cm the effect becomes important. Such fields only occur naturally inside thunderclouds. The energy distribution can be altered up to energies of $\\sim 1$~GeV. For positive fields this means positrons can outnumber the electrons over a large energy range, resulting in a positive charge excess. For negative fields electron breakdown is observed for altitudes at which the field is larger than the breakeven field and is most efficient when the field is larger than the threshold field. This electron avalanche is directed antiparallel to the electric field and can geometrically be detached from the shower. The electric field effect on the shower evolution is local in the sense that in a low field region a shower that has traversed a high field region is not much different from a shower that has not. For air showers that traverse thunderstorms this means they are generally affected only in the strong field regions of the cloud. Ground based particle detectors will not be very sensitive to these effects.  The radio signal from an air shower that travels through a strong field region can significantly change. It has been established experimentally that under thunderstorm conditions the power of the radio pulse may be much larger than anticipated. The order of magnitude of maximum electric fields that occur in thunderstorms coincide with the field strengths at which our simulations show a considerable change in particle distribution. It should be noted, however, that for the strength of the radio pulse to increase, the shower evolution does not necessarily have to change. The direct acceleration of the particles by the electric field produces radiation just like the magnetic deflection does. The effect of electric fields on the radio emission of air showers will be explored in more detail in a forthcoming paper.  Air showers can trigger electron avalanches in regions where the electric field exceeds the threshold field. These avalanches may play a role in lightning initiation. Although electron avalanches can be initiated by any seed electron of sufficient energy in a field exceeding the threshold field, a passing air shower offers the unique scenario in which a huge number of high energy electrons is injected in a very small time. The possible interaction between air showers and thunderstorms can be investigated with a hybrid array of particle detectors and radio receivers. In principle radio antennas can pick up the signal of both air showers and electrical discharges, but it is probably not possible to unambiguously detect an air shower signal behind the violent radiative background of a thunderstorm. With a combination of particle detectors and an array of radio antennas electrical activity after an air shower passage could be imaged, allowing for a study of temporal and spatial coincidences.     \\label{sec:con}  \\begin{small}  % \\vspace*{2ex} \\par \\noindent {\\em Acknowledgements.} Part of this research has been supported by Grant number VH-NG-413 of the Helmholtz Association. \\end{small}"
    },
    "0910/0910.2953_arXiv.txt": {
        "abstract": "This paper presents CCD observations of the Algol-type eclipsing binaries RZ Dra and EG Cep. The light curves have been analyzed with the PHOEBE software and Wilson-Devinney code (2003 version). A detailed photometric analysis, based on these observations, is presented for both binarity and pulsation. The results indicate semidetached systems where the secondary component fills its Roche lobe. After the subtraction of the theoretical light curve, a frequency analysis was performed in order to check for pulsations of the primary component of each system. Moreover, a period analysis was performed for each case in order to search for additional components around the eclipsing pairs. ",
        "introduction": "Eclipsing Binaries are important astrophysical objects and several of them are known to contain a pulsating component. It is interesting to study such systems, since extra information can be extracted from both their pulsation and eclipsing properties, leading to a more reliable determination of system parameters.\\\\ EG Cep and RZ Dra are referred as candidate systems for including pulsating components \\citep{S06} and therefore have been selected for observations and study. ",
        "conclusions": "The LC analysis of EG Cep and RZ Dra showed that both of them are semi-detached systems with the secondary component filling its Roche Lobe. The periodic variations of the orbital periods of these systems could be explained by adopting the existence of a tertiary component, while the steady increase of their period is probably due to the mass transfer procedure. In contrast with the O-C diagram solution, the LC analysis for both systems showed that there is not a third light contribution in the total luminosity of the system. This disagreement can be explained by taking into account the small values of mass of the third body found in each case. Finally, we could not detect any pulsation nature of the primary components of both systems. So, more accurate data (by using larger telescope and better CCD) in the future might reveal possible pulsational behaviour."
    },
    "0910/0910.0167_arXiv.txt": {
        "abstract": " ",
        "introduction": " ",
        "conclusions": "\\begin{figure} \\centering \\vspace{4.25cm} \\special{psfile=bershady_fig13a.c.eps voffset=-5 hoffset=-5 vscale=27 hscale=27} \\special{psfile=bershady_fig13b.c.eps voffset=-5 hoffset=160 vscale=27 hscale=27} \\caption{Total and Specific Grasp versus Spectral Power for a range of instruments on 4m- and 10m-class telescopes (solid and dashed lines, respectively) partially updated from Bershady et al. (2005). See text for comments on instrument efficiency.} \\end{figure}  Here we explore the sampled parameter space in spatial versus spectral information, as well as coverage versus resolution, starting with grasp and spectral power (Figure 1.13).  Recall that because reliable, consistent measurements of efficiency are unavailable for most instruments, we use grasp instead of etendue ({\\tt warning:} we really want etendue).  Note, however, that there is a factor of 6 range in the known efficiencies of instruments tabulated in this Chapter. Further note that there are two ways of viewing the specific grasp. From the perspective of staying photon-limited at high spectral resolution, high specific grasp is important. The ``flip side'' is that low specific grasp implies high angular resolution.  Figure 1.14 shows that spatial resolution is higher in NIR instruments, while spectral resolution is higher in optical instruments.  Fiber IFUs have the largest specific grasp -- reflected in the bifurcation seen in spatial resolution, i.e., fiber-fed instruments have large footprints per element ($d\\Omega$). There is a trend of decreasing specific grasp going from fiber+lenslet, lenslet, and finally to slicers. ESI has unusually large $A \\ \\times \\ d \\Omega$ for a slicer; RSS in FMS mode has the highest specific grasp overall.  \\begin{figure} \\centering \\vspace{5.5cm} \\special{psfile=bershady_fig3.14.c.eps voffset=-12 hoffset=-15 vscale=60 hscale=60 angle=0} \\caption{Spatial resolution (a) and specific grasp (b) versus spectral power for all instruments in Tables 1-4, highlighting differences between optical (filled symbols) and NIR (open symbols), as well as between different coupling methods (labeled).} \\end{figure}  Figure 1.14 and 1.15 together show that optical and near-infrared instruments trade spatial resolution for grasp; there are no high-grasp NIR instruments; the highest spectral power instruments are optical. Optical and near-infrared instruments sample comparable total information, with optical instruments sampling a broader range of trades between spatial versus spectral information.  Older NIR instruments clearly suffer from being detector-{\\it size} limited. IMACS-IFU stands out as having significantly larger number of total information elements, $N_R \\ \\times \\ N_\\Omega$, and in this sense is on-par with future-generation instruments."
    },
    "0910/0910.2738_arXiv.txt": {
        "abstract": "{Previous observations with the \\emph{Infrared Astronomical Satellite} and the \\emph{Infrared Space Observatory}, and ongoing observations with \\emph{Spitzer} and \\emph{AKARI} have led to the discovery of over 200 debris disks, based on detected mid- and far infrared excess emission, indicating warm circumstellar dust. In order to constrain the properties of these systems, e.g., to more accurately determine the dust mass, temperature and radial extent, follow-up observations in the submillimetre wavelength region are needed.} {The $\\beta$ Pictoris Moving Group is a nearby stellar association of young (${\\sim}12\\,$Myr) co-moving stars including the classical debris disk star $\\beta$ Pictoris. Due to their proximity and youth they are excellent targets when searching for submillimetre emission from cold, extended, dust components produced by collisions in Kuiper-Belt-like disks. They also allow an age independent study of debris disk properties as a function of other stellar parameters.} {We observed 7 infrared-excess stars in the $\\beta$ Pictoris Moving Group with the LABOCA bolometer array, operating at a central wavelength of 870\\,{$\\mu$}m at the 12-m submillimetre telescope APEX. The main emission at these wavelengths comes from large, cold dust grains, which constitute the main part of the total dust mass, and hence, for an optically thin case, make better estimates on the total dust mass than earlier infrared observations. Fitting the spectral energy distribution with combined optical and infrared photometry gives information on the temperature and radial extent of the disk.} {From our sample, $\\beta$~Pic, HD\\,181327, and HD\\,172555 were detected with at least 3$\\sigma$ certainty, while all others are below 2$\\sigma$ and considered non-detections. The image of $\\beta$~Pic shows an offset flux density peak located near the south-west extension of the disk, similar to the one previously found by SCUBA at the JCMT. We present SED fits for detected sources and give an upper limit on the dust mass for undetected ones.} {We find a mean fractional dust luminosity $\\bar{f}_\\mathrm{dust}=11{\\cdot}10^{-4}$ at $t \\approx 12\\,$Myr, which together with recent data at 100\\,Myr suggests an $f_\\mathrm{dust} \\propto t^{-\\alpha}$ decline of the emitting dust, with ${\\alpha}>0.8$.} ",
        "introduction": "The first large and unbiased survey of the infrared (IR) sky, conducted by the \\emph{Infrared Astronomical Satellite} (IRAS), revealed that approximately 15$\\%$ of nearby main-sequence stars have an excess of 12--100\\,$\\mu$m emission, corresponding to a luminosity at least $2\\times10^{-5}$ times higher than that expected from a pure stellar photosphere \\citep{Backman1993,Backman1987,Aumann1984}. This indicates that these stars are surrounded by warm circumstellar dust, interpreted as originating in a disk of debris left over from planet formation, which is being heated by stellar optical and ultraviolet radiation, and re-radiating at mid- and far-IR wavelengths. The short orbital lifetime of small dust particles suggests that they are being continuously replenished by collisions of larger bodies \\citep[see e.g.][]{Backman1995}. Deeper and more targeted observations by the \\emph{Infrared Space Observatory} (ISO), \\emph{Spitzer}, and \\emph{AKARI}, have revised and extended this list of debris disk candidates to encompass over 200 stars \\citep{Matthews2007,Rhee2007,DDDB}.  Although some of the most nearby debris disks have been imaged in the optical and IR, most are spatially unresolved and characterised solely on fits of the spectral energy distribution (SED) to stellar synthetic spectra and a handful of IR photometry data points. In order to better constrain the SED, and with this the temperature and radial extent of the disk, complementary submillimetre (submm) photometry is needed. Submm observations probe the Rayleigh-Jeans tail of the thermal radiation, where the measured integrated flux is roughly (sometimes with non-negligible correction) proportional to the temperature and the mass of the dust. Since the most efficient submm emitters are large and cool grains, which also dominate the mass of the disk \\citep{Zuckerman2001}, observations with instruments like the Large APEX BOlometer CAmera \\citep[LABOCA,][]{Siringo2007} on the Atacama Pathfinder EXperiment \\citep[APEX,][]{Gusten2006} telescope (operating at 870\\,{$\\mu$}m) or SCUBA \\citep{Holland1999} at the James Clerk Maxwell Telescope \\citep[JCMT,][]{Prestage1996} (operating at 850\\,{$\\mu$}m) provide good estimators for the total dust mass (assuming an optically thin disk at these wavelengths).  Another advantage of observations in the submm regime is the ability to investigate cool dust created in very extended debris disks or belts located some hundreds of AU from the star (akin to the Solar system's Kuiper Belt), compared to the inner warm dust at 1--100~AU radius studied in IR surveys. The larger disk region probed by submm, compared to IR, makes it feasible to potentially resolve nearby disks, and morphologically determine outer disk radii, dynamical interaction with unseen planets, etc.~\\citep[e.g.][]{Wyatt2008}, even though the angular resolution in the submm in general is lower than for shorter wavelengths.  Young, nearby stars, with a confirmed mid- and far-IR excess are prime targets for a submm search for cold extended disks. The $\\beta$ Pictoris Moving Group \\citep[BPMG,][]{Zuckerman2001a} is a nearby stellar association of 30 identified young ($\\sim12$\\,Myr) co-moving member stars \\citep{Zuckerman2004,Barrado1999}, harbouring the perhaps most studied of all debris disk systems, namely $\\beta$ Pictoris \\citep{Smith1984}, and also e.g., AU~Mic. In this paper we present a study of 7 main-sequence IR excess stars in the BPMG, observed at 870\\,$\\mu$m with LABOCA at APEX in Chile. The stars were selected from the sample of currently known members of BPMG fulfilling the following criteria: (1) located at $\\delta<-50\\degr$, thus, southern stars that are not accessible to JCMT, but that will be observable with ALMA, and (2) previously detected IR excess. However, several of the stars previously observed at 850\\,{$\\mu$}m, e.g.~AT~Mic and AU~Mic, had to be excluded due to time constraints. By targeting members of the same moving group we ensure that any differences in dust properties and mass reflect intrinsic, age-independent, variations, since these stars are assumed to be coeval \\citep{Mentuch2008}. This investigation is the precursor of a larger survey of 10--100~Myr old stellar associations, which will permit a statistical survey of disk properties during crucial time periods of debris disk evolution. The ultimate goal is to determine: the incidence of cold dust disks around main sequence stars; whether these disks can serve as indicators of planetary systems; the physical characteristics, e.g. chemical composition and grain size distribution, as a function of fundamental stellar parameters such as mass, metallicity and age; and the disk lifetimes. ",
        "conclusions": "We summarize the most important findings of our 870\\,$\\mu$m observations of 7 main-sequence IR-excess stars in the BPMG as follows: \\begin{itemize} \\item[\\textbullet] Out of the 7 stars observed, we made three detections. Two of the detected objects have previously never been detected at submm wavelengths. Our observations increase the frequency of detected submm disks in the BPMG to almost 17\\% (5 out of 30 stars). \\item[\\textbullet] $\\beta$~Pic showed a strong flux density peak centered on the star, with two additional weaker flux peaks, both at a PA of $\\sim36\\degr$ and radial distances of 600--700\\,AU and 2000--2500\\,AU, respectively. The former of these has a position consistent with the ``blob'' imaged by \\citet{Holland1998} at 850\\,$\\mu$m. Both may be background submm galaxies, but surprisingly show a close alignment to the observed optical and IR disk plane. Simple fitting of a stellar blackbody and a disk modified blackbody to the SED data suggests a dust temperature of 89\\,K and a low opacity law exponent $\\beta \\approx 0.7$. We calculate a minimum dust mass using the theoretically derived maximum opacity index for 1\\,mm sized icy grains and arrive at 4.8\\,$M_{\\mathrm{Moon}}$. \\item[\\textbullet] HD\\,181327 is clearly detected but at most marginally resolved with an elongation consistent with the previously imaged inclined circumstellar dust ring. SED fitting gives $T_{\\mathrm{dust}}=70$\\,K and $\\beta=0.15$, but with the 870\\,$\\mu$m flux ending up above the fit, implying the existence of a population of colder dust grains. The calculated minimum dust mass $M_\\mathrm{dust}=34\\,M_{\\mathrm{Moon}}$ is remarkably high. \\item[\\textbullet] HD\\,172555 does not show any significant submm emission at the position of the star, however, in the region between the star and its lower mass companion a very extended flux density distribution is seen. Assuming that the emission comes from dust associated with HD\\,172555 we derive the thermal equilibrium temperature of dust at the projected distances ($\\sim$1000\\,AU), which becomes 15\\,K. The integrated flux in the observed features lies far above what would be expected for a single modified blackbody disk fit to IR photometry data in the SED. Adding a second disk component with a temperature of 10--20\\,K fits nicely with the observed SED, and agrees with the temperature estimated from dust grain distances to the star. The corresponding dust mass would be 10--60\\,$M_{\\mathrm{Moon}}$ and 20--70\\,$M_{\\mathrm{Moon}}$ for the NW and SE feature, respectively. \\item[\\textbullet] We derive the fractional dust luminosity of detected sources, which is then plotted versus stellar spectral type. The mean fractional dust luminosity, $\\bar{f}_\\mathrm{dust}=11{\\cdot}10^{-4}$ measured for debris disk in the BPMG, agrees with the expected stage of collisional dust evolution for such ${\\sim}12\\,$Myr systems predicted from previous observations \\citep[e.g.][]{Su2006,Liu2004,Spangler2001} and models \\citep{Kenyon2008,Wyatt2007}. The large scatter in $f_\\mathrm{dust}$ among co-eval stars could be a sign of stochastic collisional dust production or a consequence of BMPG's evolutionary stage, perhaps just at the onset of collisional dust cascades at $t \\approx 10\\,$Myr. Comparison with data at 100\\,Myr \\citep{Greaves2009} gives $f_\\mathrm{dust} \\propto t^{-\\alpha}$ with ${\\alpha}>0.8$. \\end{itemize}"
    },
    "0910/0910.5118_arXiv.txt": {
        "abstract": "{ We investigate the number and type of pulsars that will be discovered with the low-frequency radio telescope LOFAR. We consider different search strategies for the Galaxy, for globular clusters and for other galaxies. We show that a 25-day all-sky Galactic survey can find approximately 900 new pulsars, probing the local pulsar population to a deep luminosity limit. For targets of smaller angular size such as globular clusters and galaxies many LOFAR stations can be combined coherently, to make use of the full sensitivity. Searches of nearby northern-sky globular clusters can find new low luminosity millisecond pulsars. Giant pulses from Crab-like extragalactic pulsars can be detected out to over a Mpc.  } ",
        "introduction": "Since the discovery of the first four pulsars with the Cambridge radio telescope \\citep{hbp+68}, an ongoing evolution of telescope systems has doubled the number of known radio pulsars roughly every 4 years: An evolution from a large flat receiver with a fixed beam on the sky (the original Cambridge radio telescope) to focusing dishes \\citep[Arecibo --][]{ht75b}, often steerable (Green Bank Telescope), on both hemispheres \\citep[the Parkes telescope --][]{mlc+01}, with large bandwidths and multiple simultaneously usable receivers for wider fields of view (Parkes, Arecibo). The types of pulsars discovered have changed accordingly, from slow, bright, single and nearby pulsars (the original four) to fast (young and millisecond pulsars), far-away (globular clusters) or dim pulsars, some of which are in binaries.  The next step in radio telescope evolution will be the use of large numbers of low-cost receivers that are combined interferometrically. These telescopes, the Allen Telescope Array \\citep{bow07b}, LOFAR \\citep{rot03}, MeerKAT \\citep{jon07}, ASKAP \\citep{jtb+08} and the SKA \\citep{kram04}, create new possibilities for pulsar research.  In this paper, we investigate the prospects of finding radio pulsars with LOFAR, the LOw Frequency ARray. We outline and compare strategies for targeting normal and millisecond pulsars (MSPs), both in the disk and globular clusters of our Galaxy, and in other galaxies. ",
        "conclusions": "Because of its large area and field of view, LOFAR can reveal the local population of pulsars to a very deep luminosity limit. A 25-day all-sky survey at 140\\,MHz would find 900 new pulsars, disclosing the local low-luminosity population and roughly doubling the number of pulsars known in the northern hemisphere. Millisecond pulsars in nearby globular clusters can be detected to lower flux limits than previously possible. Assuming the pulsar population in other galaxies is similar to that in ours, we can detect periodicities or giant pulses from extragalactic pulsars up to several Mpcs away."
    },
    "0910/0910.2134_arXiv.txt": {
        "abstract": "% We compare the luminosity function and rate inferred from the GBM  long bursts peak flux distribution with those inferred from the Swift and BATSE peak flux distribution. We find that  the GBM, BATSE and the Swift peak fluxes can be fitted by  the same luminosity function implying the consistency of these three samples. Using the trigger algorithm of the LAT instrument we derive important  information on the flux at 100 MeV compared to lower energy detected by the GBM. We find that the simple extension  of the synchrotron emission to high energy cannot justify the low rate of GRBs detected by LAT and for several GRBs detected by the GBM, the flux at $>100$  MeV should be suppressed.  Two bursts, GRB090217 and GRB 090202b, detected by LAT have very soft spectra in the GBM and therefore their high energy emission cannot be due to an extension of the synchrotron. ",
        "introduction": "Gamma ray bursts (GRBs) are  the most powerful events in the universe, their  total emitted energy outputs exceeding sometime $10^{54}$ ergs, owing to relativistic aberration. In GRBs in fact,  the most extreme relativistic regimes are  attained among all high energy sources of macroscopic size.  Most of the energy is radiated in gamma-rays of 100 to 1000 keV,  with tails up to the GeV domain, as detected formerly by CGRO/EGRET (Dingus 1995, Hurley et al. 1994, Gonzalez et al. 1994) and more recently and with much better detail by the AGILE/GRID and Fermi/LAT instruments (Marisaldi et al. 2009, Giuliani et al. 2008,  Giuliani et al. 2009, Abdo et al. 2009, Abdo et al. 2009, Bissaldi et al. 2009, Granot 2009).  Long GRBs  ($T_{90}>2$s) are thought to trace the history of massive star formation of the Universe and are detected all the way from locally (40 Mpc)  to the edge of the Universe ($z \\sim 8$, Salvaterra et al. 2009; Tanvir et al. 2009), so that they are ideal targets for the unbiased study of cosmological effects on the propagation of high energy photons and on the evolution of star formation.  However the number of GRBs with a measured redshift is still limited and at present we cannot derive directly the GRB luminosity function and rate evolution that are fundamental to understand the nature of these objects. We can constrain the luminosity function and rate distribution by fitting their   peak flux distributions  to those expected for a given luminosity function and GRB rate, as done for CGRO/BATSE and {\\it Swift}/BAT GRBs (Piran 1992, Cohen \\& Piran 1995, Fenimore \\& Bloom 1995, Loredo \\& Wasserman 1995, Horack \\& Hakkila 1997, Loredo \\& Wasserman 1998, Schmidt 1999, Schmidt 2001, Sethi \\& Bhargavi 2001, Guetta, Piran \\& Waxman 2005, Guetta \\& Piran 2005, 2006, 2007). Since the observed flux distribution is a convolution of these two unknown functions we must assume one and find a best fit for the other.  Here we concentrate on the $\\sim$ 250 long GRBs detected so far by the Gamma Ray Burst Monitor (GBM) onboard the Fermi GST (Table 1).  We assume that the rate of long bursts follows the star formation rate and we search for the parameters of the luminosity function. In the first part of this paper, we show that one can obtain a fully consistent fit for  the GBM, BATSE and the {\\it Swift} peak flux populations, implying that the GBM sample has properties similar to those of BATSE and {\\it Swift}.  Only 10 long bursts have been detected above  30 MeV by the LAT. Considering only the ratio between the LAT and GBM field of view (2.5 sr and more than 8 sr, respectively) and assuming comparable sensitivity of the 2 instruments in their energy ranges to the level of emission expected from GRBs, we would have expected that $\\sim 1/3$ of the GBM bursts has a $> 30$ MeV counterpart and not only $\\sim$ 5\\%, as detected.  There could be 2 effects at play: the LAT sensitivity and the fact that the intrinsic brightness of the very high energy tails of GRBs can be lower than predicted with a simple extrapolation of the soft gamma-ray spectrum or based on synchrotron and inverse Compton flux estimates (see e.g. Ando, Nakar \\& Sari 2008, who have predicted a LAT detection rate of about 20 GRB/yr and see Fan 2009).  Another possible explanation not related to instrumental or intrinsic GRB physics reasons is that in very distant GRBs with highly collimated jets, the intrinsic strong high energy emission is suppressed by the diffuse extragalactic infrared background (Gilmore,  Prada,  \\& Primack  2009).  The statistics of the LAT detection with respect to GBM detection is not dissimilar from that of the GRO instruments EGRET vs BATSE:  if the detection were only related to the FOV extension, EGRET should have detected about 50 GRBs out of the nearly 3000 detected by BATSE (BATSE covered virtually all sky, while the EGRET FOV was 30 deg in diameter).  The fact that EGRET detected only about 20 GRBs above 20 MeV, of which about 5 with the spark chamber at the higher energies (i.e. higher than 200 MeV), reflects broadly the LAT-vs-GBM statistics and indicates an intrinsic paucity of detected very high energy tails.    Similar to the LAT-detected GRBs, the EGRET-detected ones were among the brightest BATSE GRBs (Dingus et al. 1995).  Recently,  Kaneko et al. (2008) analyzed the spectral shapes of the BATSE GRBs observed by the TASC calorimeter, i.e. with  MeV emission, and found that the spectra of these bursts are quite hard  (i.e. have fairly high E$_{\\rm peak}$ or small high energy spectral index). This is similar to what we find for the bursts detected by LAT (see Fig. 1).  A proper understanding of the  link  between the GRB spectral maximum emission (100 keV  - 1 MeV) and its countepart above few MeV  is the key to get insight into the inner mechanism of power generation and into the jet formation and collimation (see Kumar \\& Barniol Duran 2009, Zhang \\& Pe'er 2009,  Zou, Fan, Piran 2009,  Li 2009), and ultimately will explain  how the GRB emission at the highest energies correlates with the fundamental GRB parameters (Amati, Frontera, Guidorzi 2009,   Ghirlanda, Nava, \\& Ghisellini 2009).  In this paper, after having verified that the properties of the GBM GRB population do not differ from those of the previous GRB missions BATSE and Swift-BAT (Section 2),   we analyze the sensitivity properties of the LAT and we estimate (following Band et al. 2009) the LAT detection rate of the high energy (100 MeV - 10 GeV) counterparts of GBM GRBs, under the hypothesis of  a simple extrapolation of the GRB spectrum to energies larger than 1 MeV.   \\begin{figure}[ht!] \\centerline{\\psfig{figure=spectrLAT.eps,height=8cm,width=9.cm}} \\caption{The spectral parameters, $\\beta$ and $E_{p}$ of the GRBs detected by LAT} \\end{figure} ",
        "conclusions": "We have shown that the simple extension of the synchrotron emission cannot justify the lack of detection of high energy emission from GRBs. The emission is suppressed at high energy by some mechanism (Fan 2009): i.e. the electrons are not accelerated to such an high energy, the high energy photons pair produce with low energy photons in the source and cannot escape the source or pair produce with external photons.  On the other hand  for the bursts detected by LAT there are two GRBs that have very soft synchrotron emission and their high energy emission cannot be explained by simple extrapolation of the low energy emission. There is evidence of another component at least in one GRB 090902b.  \\begin{figure}[ht!] \\centerline{\\psfig{figure=beta.eps,height=8cm,width=9.cm}} \\caption{ The ratio s defined in the text as a function of the high energy spectral index $\\beta$  for each burst not detected by LAT (blue cross), and detected by LAT (red strars). The measured value of 080916C is also reported (black diamond)} \\end{figure}  \\begin{figure}[ht!] \\centerline{\\psfig{figure=epeaks.eps,height=8cm,width=9.cm}} \\caption{ The ratio s defined in the text as a function of the peak energy  for each burst not detected by LAT (blue cross), and detected by LAT (red strars). The measured value of 080916C is also reported (black diamond)} \\end{figure}    \\onecolumn   \\begin{table}[t] \\caption{GBM parameters for the bursts detected by Fermi. The first column report the GRB name, the second column the duration of the GRB, the third one the GRB fluence (set to -1 when not available), the fourth the peak flux (taken in 1 sec for long bursts and set to -1 when not available), the fifth column report the method used to fit the spectra, the sixth the peak energy (set to -1 when not available), the seventh and eight the $\\alpha$ and $\\beta$ spectral indexes (set to 1000 when not available). The ninth column report the angle, $\\theta$,  from the LAT boresight, in deg (set to -1 when not available). The last column indicates LAT detection (1 = YES, 0 = NO)} \\begin{tabular}{lllllllllll} \\hline GRB & T90 &  Fluence (erg/cm2) & PF(ph/s/cm2) & Function & $E_{peak}$ & $\\alpha$ & $\\beta$ & $\\theta$ & LAT \\\\ &    sec &      $10^{-6}$erg/cm$^2$& ph/cm$^2$/sec& & keV & & & \\\\ \\hline 080810  & 122.0  & 6.90e+00  & 1.85e+00  & PL+HEC  & 313.5  & -0.91  & -1000.00  & 61  & 0 \\\\ 080812  &  15.0  & -1.00e+00  & -1.00e+00  & PL+HEC  & 140.0  &  0.17  & -1000.00  & 71  & 0 \\\\ 080816A  &  70.0  & 1.86e+01  & 3.48e+00  & PL+HEC  & 146.7  & -0.57  & -1000.00  & 55  & 0 \\\\ 080816B  &   5.0  & -1.00e+00  & 1.38e+00  & PL+HEC  & 1230.0  & -0.37  & -1000.00  & 70  & 0 \\\\ 080817A  &  70.0  & -1.00e+00  & -1.00e+00  &      *  &  -1.0  & 1000.00  & -1000.00  & 80  & 0 \\\\ 080817B  &   6.0  & 2.60e+00  & -1.00e+00  &    SPL  &  -1.0  & -1.07  & -1000.00  & 68  & 0 \\\\ 080818A  &  50.0  & 2.26e+00  & -1.00e+00  &    SPL  &  -1.0  & -1.57  & -1000.00  & 68  & 0 \\\\ 080818B  &  10.0  & 1.00e+00  & -1.00e+00  & PL+HEC  &  80.0  & -1.30  & -1000.00  & 68  & 0 \\\\ 080823  &  46.0  & 4.10e+00  & -1.00e+00  & PL+HEC  & 164.7  & -1.20  & -1000.00  & 77  & 0 \\\\ 080824  &  28.0  & 2.30e+00  & -1.00e+00  &   Band  & 100.0  & -0.40  & -2.10  & 17  & 0 \\\\ 080825C  &  22.0  & 2.40e+01  & -1.00e+00  &   Band  & 155.0  & -0.39  & -2.34  & 60  & 1 \\\\ 080830  &  45.0  & 4.60e+00  & -1.00e+00  &   Band  & 154.0  & -0.59  & -1.69  & 23  & 0 \\\\ 080904  &  22.0  & 2.25e+00  & 3.50e+00  &   Band  &  35.0  &  0.00  & -2.70  & 23  & 0 \\\\ 080905A  &   1.0  & 2.80e-01  & 6.10e+00  &    SPL  &  -1.0  & -0.96  & -1000.00  & 28  & 0 \\\\ 080905B  & 159.0  & 4.10e-02  & 2.10e-01  &    SPL  &  -1.0  & -1.75  & -1000.00  & 82  & 0 \\\\ 080905C  &  28.0  & 4.60e+00  & 4.40e+00  & PL+HEC  &  78.8  & -0.90  & -1000.00  & 108  & 0 \\\\ 080906B  &   5.0  & 1.09e+01  & 2.20e+01  &   Band  & 125.3  & -0.07  & -2.10  & 32  & 0 \\\\ 080912  &  17.0  & 3.30e+00  & 4.10e+00  &    SPL  &  -1.0  & -1.74  & -1000.00  & 56  & 0 \\\\ 080913B  & 140.0  & 2.20e+00  & -1.00e+00  & PL+HEC  & 114.0  & -0.69  & -1000.00  & 71  & 0 \\\\ 080916A  &  60.0  & 1.50e+01  & 4.50e+00  & PL+HEC  & 109.0  & -0.90  & -1000.00  & 76  & 0 \\\\ 080916C  & 100.9  & 2.40e+02  & 6.87e+00  &   Band  & 566.0  & -0.92  & -2.28  & 48  & 1 \\\\ 080920  &  85.0  & 2.40e+00  & 1.29e+00  &    SPL  &  -1.0  & -1.42  & -1000.00  & 16  & 0 \\\\ 080925  &  29.0  & 9.70e+00  & -1.00e+00  &   Band  & 120.0  & -0.53  & -2.26  & 38  & 0 \\\\ 080927  &  25.0  & 5.70e+00  & 2.00e+00  &    SPL  &  -1.0  & -1.50  & -1000.00  & 75  & 0 \\\\ 080928  &  87.0  & 1.50e+00  & -1.00e+00  &    SPL  &  -1.0  & -1.80  & -1000.00  & -1  & 0 \\\\ 081003C  &  67.0  & 5.40e+00  & -1.00e+00  &    SPL  &  -1.0  & -1.41  & -1000.00  & 48  & 0 \\\\ 081006A  &   7.0  & 7.10e-01  & -1.00e+00  &   Band  & 1135.0  & -0.77  & -1.80  & 16  & 0 \\\\ 081006B  &   9.0  & 7.30e-01  & -1.00e+00  &    SPL  &  -1.0  & -1.30  & -1000.00  & 3  & 0 \\\\ 081007A  &  12.0  & 1.20e+00  & 2.20e+00  &    SPL  &  -1.0  & -2.10  & -1000.00  & 116  & 0 \\\\ 081007B  &   0.5  & -1.00e+00  & -1.00e+00  &      *  &  -1.0  & 1000.00  & -1000.00  & -1  & 0 \\\\ 081009  &   9.0  & 3.50e+01  & -1.00e+00  &   Band  &  47.4  &  0.10  & -3.42  & 67  & 0 \\\\ 081012  &  30.0  & 3.80e+00  & 2.00e+00  & PL+HEC  & 360.0  & -0.31  & -1000.00  & 61  & 0 \\\\ 081021  &  25.0  & 5.30e+00  & 4.20e+00  &   Band  & 117.0  &  0.11  & -2.80  & 125  & 0 \\\\ 081024B  &   0.8  & 3.40e-01  & 4.20e+00  & PL+HEC  & 1583.0  & -0.70  & -1000.00  & -1  & 1 \\\\ 081024C  &  65.0  & 4.00e+00  & 1.00e+00  &   Band  &  65.0  & -0.60  & -2.50  & 78  & 0 \\\\ 081025  &  45.0  & 7.10e+00  & 4.50e+00  &   Band  & 200.0  &  0.15  & -2.05  & -1  & 0 \\\\ 081028B  &  20.0  & 2.00e+00  & 6.90e+00  & PL+HEC  &  70.0  & -0.55  & -1000.00  & 107  & 0 \\\\ 081101A  &   0.2  & 1.60e-01  & -1.00e+00  &    SPL  &  -1.0  & -1.14  & -1000.00  & -1  & 0 \\\\ 081101B  &   8.0  & 1.60e+01  & 1.03e+01  & PL+HEC  & 550.0  & -0.62  & -1000.00  & 116  & 0 \\\\  \\hline \\end{tabular} \\end{table} \\pagebreak \\newpage \\begin{table}[t] \\begin{tabular}{llllllllll} \\hline GRB name &   t90 &     fluence &       peakflux &    Function & E$_{\\rm peak}$ &  $\\alpha$ &    $\\beta$ &       $\\theta$ &  LAT\\\\ &    sec &      $10^{-6}$erg/cm$^2$& ph/cm$^2$/sec& keV & & & \\\\ \\hline 081102A  &  88.0  & 2.10e+00  & -1.00e+00  &   Band  &  72.0  &  0.44  & -2.36  & -1  & 0 \\\\ 081102B  &   2.2  & 1.12e+00  & 3.68e+00  &    SPL  &  -1.0  & -1.07  & -1000.00  & 53  & 0 \\\\ 081105B  &   0.2  & 2.28e-01  & 2.00e+01  &    SPL  &  -1.0  & -1.17  & -1000.00  & 87  & 0 \\\\ 081107  &   2.2  & 1.64e+00  & 1.10e+01  &   Band  &  65.0  &  0.25  & -2.80  & 52  & 0 \\\\ 081109A  &  45.0  & 6.53e+00  & 3.20e+00  & PL+HEC  & 240.0  & -1.28  & -1000.00  & -1  & 0 \\\\ 081110  &  20.0  & -1.00e+00  & -1.00e+00  &      *  &  -1.0  & 1000.00  & -1000.00  & 67  & 0 \\\\ 081113  &   0.5  & 1.07e+00  & 2.00e+01  &    SPL  &  -1.0  & -1.28  & -1000.00  & 60  & 0 \\\\ 081118B  &  20.0  & 1.12e-01  & 6.70e-01  &   Band  &  41.2  &  0.80  & -2.14  & 41  & 0 \\\\ 081119  &   0.8  & 4.10e-01  & 7.20e+00  &    SPL  &  -1.0  &  1.30  & -1000.00  & 86  & 0 \\\\ 081120  &  12.0  & 2.70e+00  & 5.10e+00  &   Band  &  44.0  &  0.40  & -2.18  & 84  & 0 \\\\ 081121  &  -1.0  & -1.00e+00  & -1.00e+00  &      *  &  -1.0  & 1000.00  & -1000.00  & -1  & 0 \\\\ 081122A  &  26.0  & 9.60e+00  & 3.00e+01  &   Band  & 158.6  & -0.63  & -2.24  & 21  & 0 \\\\ 081122B  &   0.3  & 7.90e-02  & 1.60e+00  &    SPL  &  -1.0  & -1.50  & -1000.00  & 52  & 0 \\\\ 081124  &  35.0  & 9.50e-02  & 6.70e-01  &   Band  &  22.8  & -0.60  & -2.83  & 86  & 0 \\\\ 081125  &  15.0  & 4.91e+01  & 2.70e+01  &   Band  & 221.0  &  0.14  & -2.34  & 126  & 0 \\\\ 081129  &  59.0  & 2.00e+01  & 1.40e+01  &   Band  & 150.0  & -0.50  & -1.84  & 118  & 0 \\\\ 081130A  &  -1.0  & -1.00e+00  & -1.00e+00  &      *  &  -1.0  & 1000.00  & -1000.00  & 90  & 0 \\\\ 081130B  &  12.0  & 1.30e+00  & 1.80e+00  & PL+HEC  & 152.0  & -0.77  & -1000.00  & 66  & 0 \\\\ 081204B  &   0.3  & 4.88e-01  & 1.63e+01  &    SPL  &  -1.0  & -1.18  & -1000.00  & 46  & 0 \\\\ 081204C  &   4.7  & 1.48e+00  & 7.20e+00  &    SPL  &  -1.0  & -1.40  & -1000.00  & 56  & 0 \\\\ 081206A  &  24.0  & 4.00e+00  & 2.40e+00  &   Band  & 151.0  &  0.13  & -2.20  & 102  & 0 \\\\ 081206B  &  10.0  & -1.00e+00  & -1.00e+00  &      *  &  -1.0  & 1000.00  & -1000.00  & 82  & 0 \\\\ 081206C  &  20.0  & 1.19e+00  & 7.60e-01  &    SPL  &  -1.0  & -1.35  & -1000.00  & 71  & 0 \\\\ 081207  & 153.0  & 1.06e+02  & -1.00e+00  &   Band  & 639.0  & -0.65  & -2.41  & 56  & 0 \\\\ 081209  &   0.4  & 5.90e-01  & 7.80e+00  &   Band  & 808.0  & -0.50  & -2.00  & 109  & 0 \\\\ 081213  &   0.0  & -1.00e+00  & -1.00e+00  &      *  &  -1.0  & 1000.00  & -1000.00  & 51  & 0 \\\\ 081215A  &   7.7  & 5.44e+01  & 6.89e+01  &   Band  & 304.0  & -0.58  & -2.07  & 86  & 1 \\\\ 081215B  &  90.0  & 2.80e+00  & -1.00e+00  & PL+HEC  & 139.0  & -0.14  & -1000.00  & 112  & 0 \\\\ 081216  &   1.0  & 3.60e+00  & 5.50e+01  &   Band  & 1235.0  & -0.70  & -2.17  & 95  & 0 \\\\ 081217  &  39.0  & 1.00e+01  & 4.00e+00  &   Band  & 167.0  & -0.61  & -2.70  & 54  & 0 \\\\ 081221  &  40.0  & 3.70e+01  & 3.30e+01  &   Band  &  77.0  & -0.42  & -2.91  & 78  & 0 \\\\ 081222  &  30.0  & 1.35e+01  & 1.48e+01  &   Band  & 134.0  & -0.55  & -2.10  & 50  & 0 \\\\ 081223  &   0.9  & 1.20e+00  & 2.20e+01  & PL+HEC  & 280.0  & -0.63  & -1000.00  & 28  & 0 \\\\ 081224  &  50.0  & -1.00e+00  & -1.00e+00  &      *  &  -1.0  & 1000.00  & -1000.00  & 16  & 0 \\\\ 081225  &  42.0  & 2.45e+00  & 6.00e-01  &    SPL  &  -1.0  & -1.51  & -1000.00  & 55  & 0 \\\\ 081226A  &   1.7  & 2.10e-01  & 3.20e+00  &    SPL  &  -1.0  &  1.17  & -1000.00  & 110  & 0 \\\\ 081226B  &   0.4  & 6.10e-01  & 1.77e+01  &   Band  & 300.0  & -0.20  & -1.82  & 22  & 0 \\\\ 081226C  &  60.0  & 2.32e+00  & 4.50e+00  & PL+HEC  &  82.0  & -1.04  & -1000.00  & 54  & 0 \\\\ 081229  &   0.5  & 8.70e-01  & 1.07e+01  &   Band  & 585.0  & -0.27  & -2.00  & 44  & 0 \\\\ \\hline \\end{tabular} \\end{table} \\pagebreak \\newpage \\begin{table}[t] \\begin{tabular}{llllllllll} \\hline GRB name &   t90 &     fluence &       peakflux &    Function & E$_{\\rm peak}$ &  $\\alpha$ &    $\\beta$ &       $\\theta$ &  LAT\\\\ &    sec &      $10^{-6}$erg/cm$^2$& ph/cm$^2$/sec& keV & & & \\\\ \\hline 081231  &  29.0  & 1.20e+01  & 1.53e+00  &   Band  & 152.3  & -0.80  & -2.03  & 20  & 0 \\\\ 090107B  &  24.1  & 1.75e+00  & 3.68e+00  & PL+HEC  & 106.1  & -0.68  & -1000.00  & -1  & 0 \\\\ 090108A  &   0.9  & 1.28e+00  & 3.97e+01  &   Band  & 104.8  & -0.47  & -1.97  & 60  & 0 \\\\ 090108B  &   0.8  & 7.90e-01  & 1.90e+01  &    SPL  &  -1.0  & -0.99  & -1000.00  & 72  & 0 \\\\ 090109  &   5.0  & 1.21e+00  & 2.76e+00  &    SPL  &  -1.0  & -1.50  & -1000.00  & 62  & 0 \\\\ 090112A  &  65.0  & 5.20e+00  & 7.00e+00  &   Band  & 150.0  & -0.94  & -2.01  & 4  & 0 \\\\ 090112B  &  12.0  & 5.40e+00  & 1.40e+01  &   Band  & 139.0  & -0.75  & -2.43  & 95  & 0 \\\\ 090117A  &  21.0  & 1.80e+00  & 9.60e+00  &   Band  &  25.0  & -0.40  & -2.50  & 51  & 0 \\\\ 090117B  &  27.0  & 2.10e+00  & 4.60e+00  &    SPL  &  -1.0  & -1.55  & -1000.00  & 49  & 0 \\\\ 090117C  &  86.0  & 1.10e+01  & 4.20e+00  &   Band  & 247.0  & -1.00  & -2.10  & 54  & 0 \\\\ 090126B  &  10.8  & 1.25e+00  & 4.90e+00  & PL+HEC  &  47.5  & -0.99  & -1000.00  & 18  & 0 \\\\ 090126C  &  -1.0  & -1.00e+00  & -1.00e+00  &      *  &  -1.0  & 1000.00  & -1000.00  & 68  & 0 \\\\ 090129  &  17.2  & 5.60e+00  & 8.00e+00  &   Band  & 123.2  & -1.39  & -1.98  & 22  & 0 \\\\ 090131  &  36.4  & 2.23e+01  & 4.79e+01  &   Band  &  58.4  & -1.27  & -2.26  & 40  & 0 \\\\ 090202  &  66.0  & 8.65e+00  & 7.77e+00  & PL+HEC  & 570.0  & -1.31  & -1000.00  & 55  & 0 \\\\ 090206  &   0.8  & 1.04e+00  & 1.90e+01  & PL+HEC  & 710.0  & -0.65  & -1000.00  & 72  & 0 \\\\ 090207  &  10.0  & 4.01e+00  & 1.88e+00  &    SPL  &  -1.0  & -1.59  & -1000.00  & 45  & 0 \\\\ 090213  &  -1.0  & -1.00e+00  & -1.00e+00  &      *  &  -1.0  & 1000.00  & -1000.00  & 17  & 0 \\\\ 090217  &  32.8  & 3.08e+01  & 1.12e+01  &   Band  & 610.0  & -0.85  & -2.86  & 34  & 1 \\\\ 090219  &   0.5  & 8.00e-01  & 7.70e+00  &    SPL  &  -1.0  & -1.43  & -1000.00  & 137  & 0 \\\\ 090222  &  18.0  & 2.19e+00  & 1.10e+00  &   Band  & 147.9  & -0.97  & -2.56  & 80  & 0 \\\\ 090227A  &  50.0  & 9.00e+00  & 4.57e+00  &   Band  & 1357.0  & -0.92  & -3.60  & 21  & 0 \\\\ 090227B  &   0.9  & 8.70e+00  & 3.46e+01  &   Band  & 2255.0  & -0.53  & -3.04  & 72  & 0 \\\\ 090228A  &   0.8  & 6.10e+00  & 1.33e+02  &   Band  & 849.0  & -0.35  & -2.98  & 16  & 0 \\\\ 090228B  &   7.2  & 9.96e-01  & 2.53e+00  & PL+HEC  & 147.8  & -0.70  & -1000.00  & 20  & 0 \\\\ 090301B  &  28.0  & 2.69e+00  & 4.40e+00  &   Band  & 427.0  & -0.98  & -1.93  & 56  & 0 \\\\ 090303  &  -1.0  & -1.00e+00  & -1.00e+00  &      *  &  -1.0  & 1000.00  & -1000.00  & -1  & 0 \\\\ 090304  &  -1.0  & -1.00e+00  & -1.00e+00  &      *  &  -1.0  & 1000.00  & -1000.00  & 31  & 0 \\\\ 090305B  &   2.0  & 2.70e+00  & 1.10e+01  &   Band  & 770.0  & -0.50  & -1.90  & 40  & 0 \\\\ 090306C  &  38.8  & 9.00e-01  & 2.40e+00  &   Band  &  87.0  & -0.32  & -2.28  & 14  & 0 \\\\ 090307B  &  30.0  & 1.70e+00  & 1.80e+00  & PL+HEC  & 212.0  & -0.70  & -1000.00  & 83  & 0 \\\\ 090308B  &   2.1  & 3.46e+00  & 1.42e+01  & PL+HEC  & 710.3  & -0.54  & -1000.00  & 50  & 0 \\\\ 090309B  &  60.0  & 4.70e+00  & 4.43e+00  & PL+HEC  & 197.0  & -1.52  & -1000.00  & 26  & 0 \\\\ 090310  & 125.2  & 2.15e+00  & 4.40e+00  & PL+HEC  & 279.0  & -0.65  & -1000.00  & 77  & 0 \\\\ 090319  &  67.7  & 7.47e+00  & 3.85e+00  & PL+HEC  & 187.3  &  0.90  & -1000.00  & 27  & 0 \\\\ 090320A  &  10.0  & -1.00e+00  & -1.00e+00  &      *  &  -1.0  & 1000.00  & -1000.00  & 60  & 0 \\\\ 090320B  &  52.0  & 1.10e+00  & 1.20e-01  & PL+HEC  &  72.0  & -1.10  & -1000.00  & 101  & 0 \\\\ 090320C  &   4.0  & -1.00e+00  & -1.00e+00  &      *  &  -1.0  & 1000.00  & -1000.00  & 40  & 0 \\\\ 090323  &  70.0  & 1.00e+02  & 1.23e+01  & PL+HEC  & 697.0  & -0.89  & -1000.00  & -1  & 1 \\\\ \\hline \\end{tabular} \\end{table} \\pagebreak \\newpage \\begin{table}[t] \\begin{tabular}{llllllllll} \\hline GRB name &   t90 &     fluence &       peakflux &    Function & E$_{\\rm peak}$ &  $\\alpha$ &    $\\beta$ &       $\\theta$ &  LAT\\\\ &    sec &      $10^{-6}$erg/cm$^2$& ph/cm$^2$/sec& keV & & & \\\\ \\hline 090326  &  11.2  & 8.60e-01  & -1.00e+00  & PL+HEC  &  75.0  & -0.86  & -1000.00  & 103  & 0 \\\\ 090327  &  24.0  & 3.00e+00  & 3.50e+00  &   Band  &  89.7  & -0.39  & -2.90  & 66  & 0 \\\\ 090328A  & 100.0  & 8.09e+01  & 1.85e+01  &   Band  & 653.0  & -0.93  & -2.20  & -1  & 1 \\\\ 090328B  &   0.3  & 9.61e-01  & 2.98e+01  &   Band  & 1967.0  & -0.92  & -2.48  & 74  & 0 \\\\ 090330  &  80.0  & 1.14e+01  & 6.80e+00  &   Band  & 246.0  & -0.99  & -2.68  & 50  & 0 \\\\ 090331  &  -1.0  & -1.00e+00  & -1.00e+00  &      *  &  -1.0  & 1000.00  & -1000.00  & 40  & 0 \\\\ 090403  &  16.0  & -1.00e+00  & -1.00e+00  &      *  &  -1.0  & 1000.00  & -1000.00  & 42  & 0 \\\\ 090405  &   1.2  & -1.00e+00  & -1.00e+00  &      *  &  -1.0  & 1000.00  & -1000.00  & 70  & 0 \\\\ 090409  &  20.0  & 6.14e-01  & 1.36e+00  & PL+HEC  & 137.0  &  1.20  & -1000.00  & 90  & 0 \\\\ 090411A  &  24.6  & 8.60e+00  & 3.25e+00  &   Band  & 141.0  & -0.88  & -1.82  & 59  & 0 \\\\ 090411B  &  18.7  & 8.00e+00  & 7.40e+00  &   Band  & 189.0  & -0.80  & -2.00  & 111  & 0 \\\\ 090412  &   0.5  & -1.00e+00  & -1.00e+00  &      *  &  -1.0  & 1000.00  & -1000.00  & 71  & 0 \\\\ 090418C  &   0.6  & 6.00e-01  & 8.50e+00  & PL+HEC  & 1000.0  & -0.94  & -1000.00  & 58  & 0 \\\\ 090422  &  10.0  & 1.00e+00  & 7.80e+00  &    SPL  &  -1.0  &  1.81  & -1000.00  & 29  & 0 \\\\ 090423  &  12.0  & 1.10e+00  & 3.30e+00  & PL+HEC  &  82.0  & -0.77  & -1000.00  & 75.6  & 0 \\\\ 090424  &  52.0  & 5.20e+01  & 1.37e+02  &   Band  & 177.0  &  0.90  & -2.90  & 71  & 0 \\\\ 090425  &  72.0  & 1.30e+01  & 1.40e+01  &   Band  &  69.0  & -1.29  & -2.03  & 105  & 0 \\\\ 090426B  &   3.8  & 5.20e-01  & -1.00e+00  &    SPL  &  -1.0  & -1.60  & -1000.00  & 56  & 0 \\\\ 090426C  &  12.0  & 3.10e+00  & 6.80e+00  &   Band  & 295.0  & -1.29  & -1.98  & 69  & 0 \\\\ 090427B  &   7.0  & 8.00e-01  & -1.00e+00  &    SPL  &  -1.0  & -1.10  & -1000.00  & 14  & 0 \\\\ 090427C  &  12.5  & 1.60e+00  & -1.00e+00  & PL+HEC  &  75.0  &  0.35  & -1000.00  & 81  & 0 \\\\ 090428A  &   8.0  & 9.90e-01  & 1.23e+01  &   Band  &  85.0  & -0.40  & -2.70  & 96  & 0 \\\\ 090428B  &  30.0  & 5.20e+00  & 1.01e+01  &   Band  &  53.0  & -1.81  & -2.17  & 101  & 0 \\\\ 090429C  &  13.0  & 3.70e+00  & 6.70e+00  &    SPL  &  -1.0  & -1.43  & -1000.00  & 112  & 0 \\\\ 090429D  &  11.0  & 1.60e+00  & 8.60e-01  &    SPL  & 223.0  & -0.87  & -1000.00  & 33  & 0 \\\\ 090502  &  66.2  & 3.50e-02  & 6.20e+00  & PL+HEC  &  63.2  & -1.10  & -1000.00  & 77  & 0 \\\\ 090509  & 295.0  & 8.40e+00  & 3.10e+00  & PL+HEC  & 343.0  & -0.90  & -1000.00  & 75  & 0 \\\\ 090510A  &   1.4  & 3.00e+01  & 8.00e+01  &   Band  & 4400.0  & -0.80  & -2.60  & -1  & 1 \\\\ 090510B  &   7.0  & -1.00e+00  & -1.00e+00  &      *  &  -1.0  & 1000.00  & -1000.00  & 100  & 0 \\\\ 090511  &  14.0  & 1.80e+00  & 2.50e+00  & PL+HEC  & 391.0  & -0.95  & -1000.00  & 67  & 0 \\\\ 090513A  &  23.0  & 6.80e+00  & 2.70e+00  & PL+HEC  & 850.0  & -0.90  & -1000.00  & 89  & 0 \\\\ 090513B  &  -1.0  & -1.00e+00  & -1.00e+00  &      *  &  -1.0  & 1000.00  & -1000.00  & 119  & 0 \\\\ 090514  &  49.0  & 8.10e+00  & 7.60e+00  &    SPL  &  -1.0  & -1.92  & -1000.00  & 19  & 0 \\\\ 090516A  & 350.0  & 2.30e+01  & 5.30e+00  &   Band  &  51.4  & -1.03  & -2.10  & 20  & 0 \\\\ 090516B  & 350.0  & 3.00e+01  & 4.00e+00  & PL+HEC  & 327.0  & -1.01  & -1000.00  & 45  & 0 \\\\ 090516C  &  15.0  & 4.00e+00  & 7.70e+00  &   Band  &  38.0  & -0.44  & -1.81  & 69  & 0 \\\\ 090518A  &   9.0  & 1.60e+00  & 4.70e+00  &    SPL  &  -1.0  & -1.59  & -1000.00  & 53  & 0 \\\\ 090518B  &  12.0  & 2.20e+00  & 5.60e+00  & PL+HEC  & 127.0  & -0.74  & -1000.00  & 90  & 0 \\\\ 090519B  &  87.0  & 1.40e+00  & 5.02e+00  &    SPL  &  -1.0  & -1.63  & -1000.00  & 18  & 0 \\\\ 090520B  &   1.5  & 4.50e-01  & 4.10e+00  &    SPL  &  -1.0  & -1.40  & -1000.00  & 10  & 0 \\\\ 090520C  &   4.9  & 3.54e+00  & 4.47e+00  &   Band  & 204.2  & -0.73  & -1.96  & 71  & 0 \\\\ 090520D  &  12.0  & 4.00e+00  & 4.10e+00  &   Band  &  46.3  & -0.99  & -3.25  & 66  & 0 \\\\ 090522  &  22.0  & 1.20e+00  & 3.50e+00  & PL+HEC  &  75.8  & -1.03  & -1000.00  & 53  & 0 \\\\ 090524  &  72.0  & 1.85e+01  & 1.41e+01  &   Band  &  82.6  & -1.00  & -2.30  & 63  & 0 \\\\ 090528A  &  68.0  & 9.30e+00  & 7.60e+00  & PL+HEC  &  99.0  & -1.70  & -1000.00  & 81  & 0 \\\\ 090528B  & 102.0  & 4.65e+01  & 1.47e+01  &   Band  & 172.0  & -1.10  & -2.30  & 65  & 0 \\\\ \\hline \\end{tabular} \\end{table} \\pagebreak \\newpage \\begin{table}[t] \\begin{tabular}{llllllllll} \\hline GRB name &   t90 &     fluence &       peakflux &    Function & E$_{\\rm peak}$ &  $\\alpha$ &    $\\beta$ &       $\\theta$ &  LAT\\\\ &    sec &      $10^{-6}$erg/cm$^2$& ph/cm$^2$/sec& keV & & & \\\\ \\hline 090529B  &   5.1  & 3.40e-01  & 4.10e+00  &   Band  & 142.0  & -0.70  & -2.00  & 36  & 0 \\\\ 090529C  &  10.4  & 3.10e+00  & 2.50e+01  &   Band  & 188.0  & -0.84  & -2.10  & 69  & 0 \\\\ 090530B  & 194.0  & 5.90e+01  & 1.08e+01  &   Band  &  67.0  & -0.71  & -2.42  & 84  & 0 \\\\ 090531B  &   2.0  & 6.20e-01  & 1.49e+00  &   Band  & 2166.0  & -0.71  & -2.47  & 25  & 0 \\\\ 090602  &  16.0  & 5.70e+00  & 3.62e+00  & PL+HEC  & 503.0  & -0.56  & -1000.00  & 112  & 0 \\\\ 090606  &  60.0  & 3.19e+00  & 2.41e+00  &    SPL  &  -1.0  & -1.63  & -1000.00  & 128  & 0 \\\\ 090608  &  61.0  & 3.20e+00  & 2.70e+00  &    SPL  &  -1.0  & -1.83  & -1000.00  & 93  & 0 \\\\ 090610A  &   6.5  & 7.32e-01  & 9.40e-01  &    SPL  &  -1.0  & -1.30  & -1000.00  & 70  & 0 \\\\ 090610B  & 202.5  & 4.13e+00  & 1.54e+00  & PL+HEC  & 104.9  & -0.46  & -1000.00  & 91  & 0 \\\\ 090610C  &  18.1  & 8.54e-01  & 1.12e+00  &    SPL  &  -1.0  & -1.62  & -1000.00  & 104  & 0 \\\\ 090612  &  58.0  & 2.37e+00  & 1.63e+00  &   Band  & 357.0  & -0.60  & -1.90  & 56  & 0 \\\\ 090616  &   2.7  & 2.23e-01  & 2.08e+00  &    SPL  &  -1.0  & -1.27  & -1000.00  & 68  & 0 \\\\ 090617  &   0.4  & 4.68e-01  & 1.00e+01  &   Band  & 684.0  & -0.45  & -2.00  & 45  & 0 \\\\ 090618  & 155.0  & 2.70e+02  & 7.34e+01  &   Band  & 155.5  & -1.26  & -2.50  & 133  & 0 \\\\ 090620  &  16.5  & 6.60e+00  & 7.00e+00  &   Band  & 156.0  & -0.40  & -2.44  & 60  & 0 \\\\ 090621A  & 294.0  & 4.40e+00  & 1.92e+00  &   Band  &  56.0  & -1.10  & -2.12  & 12  & 0 \\\\ 090621B  &   0.1  & 3.71e-01  & 6.40e+00  &   Band  & 321.6  & -0.13  & -1.57  & 108  & 0 \\\\ 090621C  &  59.9  & 1.80e+00  & 2.29e+00  & PL+HEC  & 148.0  & -1.40  & -1000.00  & 52  & 0 \\\\ 090621D  &  39.9  & 1.34e+00  & 1.74e+00  &    SPL  &  -1.0  & -1.66  & -1000.00  & 79  & 0 \\\\ 090623  &  72.2  & 9.60e+00  & 3.30e+00  &   Band  & 428.0  & -0.69  & -2.30  & 73  & 0 \\\\ 090625A  &  51.0  & 8.80e-01  & 5.00e-01  & PL+HEC  & 198.0  & -0.60  & -1000.00  & 13  & 0 \\\\ 090625B  &  13.6  & 1.04e+00  & 1.87e+00  &   Band  & 100.0  & -0.40  & -2.00  & 125  & 0 \\\\ 090626  &  70.0  & 3.50e+01  & 1.79e+01  &   Band  & 175.0  & -1.29  & -1.98  & 15  & 1 \\\\ 090630  &   5.1  & 5.10e-01  & 2.78e+00  &   Band  &  71.0  & -1.50  & -2.30  & 75  & 0 \\\\ 090701  &  12.0  & 4.50e-01  & 2.10e+00  &    SPL  &  -1.0  &  1.84  & -1000.00  & 13  & 0 \\\\ 090703  &   9.0  & 6.80e-01  & 1.00e+00  &    SPL  &  -1.0  & -1.72  & -1000.00  & 25  & 0 \\\\ 090704  &  70.0  & 5.80e+00  & 1.20e+00  & PL+HEC  & 233.7  & -1.13  & -1000.00  & 77  & 0 \\\\ 090706  & 100.0  & 1.50e+00  & 1.24e+00  &    SPL  &  -1.0  & -2.16  & -1000.00  & 20  & 0 \\\\ 090708  &  18.0  & 4.00e-01  & 1.00e+00  & PL+HEC  &  47.5  & -1.29  & -1000.00  & 55  & 0 \\\\ 090709B  &  32.0  & 1.30e+00  & 2.00e+00  & PL+HEC  & 130.0  & -1.01  & -1000.00  & 35  & 0 \\\\ 090711  & 100.0  & 1.17e+01  & 4.20e+00  & PL+HEC  & 210.0  & -1.30  & -1000.00  & 13  & 0 \\\\ 090712  &  72.0  & 4.20e+00  & 6.30e-01  & PL+HEC  & 505.0  & -0.68  & -1000.00  & 33  & 0 \\\\ 090713  & 113.0  & 3.70e+00  & 1.60e+00  & PL+HEC  &  99.0  & -0.34  & -1000.00  & 63  & 0 \\\\ 090717A  &  70.0  & 4.50e-01  & 7.80e+00  &   Band  & 120.0  & -0.88  & -2.33  & 70  & 0 \\\\ 090717B  &   0.9  & 4.83e-01  & 3.91e+00  &    SPL  &  -1.0  & -1.02  & -1000.00  & 35  & 0 \\\\ 090718A  &  -1.0  & -1.00e+00  & -1.00e+00  &      *  &  -1.0  & 1000.00  & -1000.00  & 51  & 0 \\\\ 090718B  &  28.0  & 2.52e+01  & 3.20e+01  &   Band  & 184.0  & -1.18  & -2.59  & 76  & 0 \\\\ \\hline \\end{tabular} \\end{table} \\pagebreak \\newpage \\begin{table}[t] \\begin{tabular}{llllllllll} \\hline GRB name &   t90 &     fluence &       peakflux &    Function & E$_{\\rm peak}$ &  $\\alpha$ &    $\\beta$ &       $\\theta$ &  LAT\\\\ &    sec &      $10^{-6}$erg/cm$^2$& ph/cm$^2$/sec& keV & & & \\\\ \\hline 090719  &  16.0  & 4.83e+01  & 3.78e+01  &   Band  & 254.0  & -0.68  & -2.92  & 88  & 0 \\\\ 090720A  &   7.0  & 2.90e+00  & 1.09e+01  & PL+HEC  & 117.5  & -0.75  & -1000.00  & 113  & 0 \\\\ 090720B  &  20.0  & 1.06e+01  & 1.09e+01  &   Band  & 924.0  & -1.00  & -2.43  & 56  & 0 \\\\ 090802A  &   0.1  & 6.50e-01  & 6.12e+01  &   Band  & 283.0  & -0.42  & -2.40  & 123  & 0 \\\\ 090802B  &  -1.0  & -1.00e+00  & -1.00e+00  &      *  &  -1.0  & 1000.00  & -1000.00  & 104  & 0 \\\\ 090807B  &   3.0  & 1.02e+00  & 1.09e+01  &   Band  &  37.0  & -0.60  & -2.40  & 45  & 0 \\\\ 090809B  &  15.0  & 2.26e+01  & 2.36e+01  &   Band  & 198.0  & -0.85  & -2.02  & 81  & 0 \\\\ 090813  &   9.0  & 3.50e+00  & 1.44e+01  &   Band  &  95.0  & -1.25  & -2.00  & 35.3  & 0 \\\\ 090814C  &   0.2  & 6.60e-01  & 9.10e+00  & PL+HEC  & 790.0  & -0.39  & -1000.00  & 61  & 0 \\\\ 090815A  & 200.0  & 3.40e+00  & 1.90e+00  &    SPL  &  -1.0  & -1.50  & -1000.00  & 87  & 0 \\\\ 090815B  &  30.0  & 5.05e+00  & 1.44e+01  &   Band  &  18.6  & -1.82  & -2.70  & 82  & 0 \\\\ 090817  & 220.0  & 7.30e+00  & 3.80e+00  &   Band  & 115.0  & -1.10  & -2.20  & 82  & 0 \\\\ 090820A  &  60.0  & 6.60e+01  & 5.80e+01  &   Band  & 215.0  & -0.69  & -2.61  & 108  & 0 \\\\ 090820B  &  11.2  & 1.16e+00  & 6.10e+00  & PL+HEC  &  38.8  & -1.44  & -1000.00  & 32  & 0 \\\\ 090826  &   8.5  & 1.26e+00  & 3.28e+00  & PL+HEC  & 172.0  & -0.96  & -1000.00  & 35  & 0 \\\\ 090828  & 100.0  & 2.52e+01  & 1.62e+01  &   Band  & 136.5  & -1.23  & -2.12  & 95  & 0 \\\\ 090829A  &  85.0  & 1.02e+02  & 5.15e+01  &   Band  & 183.0  & -1.44  & -2.10  & 47  & 0 \\\\ 090829B  & 100.0  & 6.40e+00  & 3.20e+00  &   Band  & 143.0  & -0.70  & -2.40  & 42  & 0 \\\\ 090831  &  69.1  & 1.66e+01  & 9.40e+00  &   Band  & 243.8  & -1.52  & -1.96  & 107  & 0 \\\\ 090902A  &   1.2  & 2.11e+00  & 1.14e+01  &   Band  & 388.0  &  0.30  & -2.05  & 82  & 0 \\\\ 090902B  &  21.0  & 3.74e+02  & 4.61e+01  &   Band  & 798.0  & -0.61  & -3.87  & 52  & 1 \\\\ 090904B  &  71.0  & 2.44e+01  & 9.80e+00  &   Band  & 106.3  & -1.26  & -2.18  & 113  & 0 \\\\ 090910  &  62.0  & 9.20e+00  & 2.30e+00  &   Band  & 274.8  & -0.90  & -2.00  & 107  & 0 \\\\ 090922A  &  92.0  & 1.14e+01  & 1.56e+01  &   Band  & 139.3  & -0.77  & -2.28  & 19  & 0 \\\\ 090925  &  50.0  & 9.46e+00  & 4.20e+00  &   Band  & 156.0  & -0.60  & -1.91  & 116  & 0 \\\\ 090926A  &  20.0  & 1.45e+02  & 8.08e+01  &   Band  & 268.0  & -0.69  & -2.34  & 52  & 1 \\\\ 090926B  &  81.0  & 8.70e+00  & -1.00e+00  & PL+HEC  &  91.0  & -0.13  & -1000.00  & 100  & 0 \\\\ 090927  &   2.0  & 6.10e-01  & 7.20e+00  &    SPL  &  -1.0  & -1.47  & -1000.00  & 85  & 0 \\\\ 090929A  &   8.5  & 1.06e+01  & 1.09e+01  & PL+HEC  & 610.9  & -0.52  & -1000.00  & 122  & 0 \\\\ 091003A  &  21.1  & 3.76e+01  & 3.18e+01  &   Band  & 486.2  & -1.13  & -2.64  & 13  & 1 \\\\ \\label{t:fit} \\end{tabular} \\end{table}    \\begin{table}[t] \\caption{ Best fit parameters $Rate(z=0)$ , $L^*$, $\\alpha$ and $\\beta$ and their 1-$\\sigma$ confidence levels. For each fit we report the  $\\chi^2$ values corresponding to the best fit ($\\chi^2_{\\rm b.f.}$). Also shown are the parameter for a LF, LFb, that fit quite well all the samples, the $\\chi^2\\sim 1.3 $ for all the instruments } \\begin{tabular}{|c|c|c|c|c|c|} \\hline sample & Rate(z=0)  & $L^*$ & $\\alpha$ & $\\beta$ &$\\chi^2$\\\\ & $Gpc^{-3} yr^{-1}$ & $10^{51}$ erg/sec &  & &   \\\\ \\hline GBM & $0.5^{+0.3}_{-0.2}$ & $5.5^{+1.5}_{-2}$ & $0.3^{+0.1}_{-0.5}$ &$2.3^{+0.6}_{-0.3}$ & 1.1  \\\\ BATSE & $1.0^{+0.2}_{-0.4}$ & $4_{-1.5}^{+2}$ & $0.1_{-0.05}^{+0.3}$ &$2.6_{-0.5}^{+0.9}$ & 1.1 \\\\ {\\it Swift} & $0.6^{+0.3}_{-0.1}$ & $3.3_{-0.5}^{+2.5}$ & $0.1_{-0.05}^{+0.3}$ &$2.7_{-0.4}^{+1}$ & 0.95 \\\\ LFb &   $0.8$  & $5 $ & $0.1$ & $2.7$ & 1.3 \\\\ \\hline \\end{tabular} \\end{table}"
    },
    "0910/0910.0241_arXiv.txt": {
        "abstract": "We generalize the Swiss-cheese cosmologies so as to include nonzero linear momenta of the associated boundary surfaces. The evolution of mass scales in these generalized cosmologies is studied for a variety of models for the background without having to specify any details within the local inhomogeneities. We find that the final effective gravitational mass and size of the evolving inhomogeneities depends on their linear momenta but these properties are essentially unaffected by the details of the background model. ",
        "introduction": "The Swiss-cheese models give us (noninteracting) inhomogeneities in a cosmological setting that are exact solutions to Einstein's equations. As a result, the models have become a standard construction \\cite{forms} and are very widely studied \\cite{2}. The classical Einstein-Straus vacuole model, which requires a comoving boundary surface, is unstable. Many subsequent Swiss-cheese models have also assumed that the associated boundary surfaces remain comoving, but it is well known that this need not be the case. Some general studies of this issue go back many years \\cite{lake}. Some specific examples of noncomoving boundary surfaces include the study of Vaidya-type inhomogeneities \\cite{fayos} and the evolution of ``density waves\" in the Lema\\^{\\i}tre-Tolman model \\cite{lt}. Here we examine the role of the linear momentum of a boundary surface in a Robertson-Walker background. Within the context of these generalized models, we study the evolution of the effective gravitational mass and size of the inhomogeneities for a variety of well-known background models. We find that whereas the momentum plays a central role, the details of the background model are relatively unimportant. ",
        "conclusions": "We have introduced a new generalized Swiss-cheese model which does not assume \\emph{a priori} that the associated boundary surfaces are comoving. In order to quantify evolving inhomogeneities, we have considered geodesic boundaries characterized by their linear momentum $\\mathcal{D}$. For the size of the inhomogeneities we are interested in, the physical values of $\\mathcal{D}/c$ are $\\sim10^{-5}$. For a given linear momentum, we have found that the inhomogeneities grow almost independently of the background model (with the inclusion of the radiation density parameter $\\Omega_{R_0}$). As shown in Fig.\\ref{rsigmabar} these inhomogeneities are almost at their full size by the decoupling ($z\\sim 1100$). For a redshift of $z \\lesssim 2$, corresponding to high redshift supernovae, the inhomogeneities considered here are growing very slowly as is shown in see Fig.\\ref{MSMLCDM}."
    },
    "0910/0910.0077_arXiv.txt": {
        "abstract": "{ We present a large sample of candidate galaxies at $z\\approx 7$ -- 10, selected in the Hubble Ultra Deep Field using the new observations of the Wide Field Camera 3 that was recently installed to Hubble Space Telescope. Our sample is composed of 20 $z_{850}$-dropouts (four new discoveries), 15 $Y_{105}$-dropouts (nine new discoveries) and 20 $J_{125}$-dropouts (all new discoveries). The surface densities of the $z_{850}$-dropouts are close to what predicted by earlier studies, however, those of the $Y_{105}$- and $J_{125}$-dropouts are quite unexpected. While no $Y_{105}$- or $J_{125}$-dropouts have been found at $AB\\leq 28.0$~mag, their surface densities seem to increase sharply at fainter levels. While some of these candidates seem to be close to foreground galaxies and thus could possibly be gravitationally lensed, the overall surface densities after excluding such cases are still much higher than what would be expected if the luminosity function does not evolve from $z\\sim 7$ to 10. Motivated by such steep increases, we tentatively propose a set of Schechter function parameters to describe the luminosity functions at $z\\approx 8$ and 10. As compared to their counterpart at $z\\approx 7$, here $L^*$ decreases by a factor of $\\sim 6.5$ and $\\Phi^*$ increases by a factor of 17--90. Although such parameters are not yet demanded by the existing observations, they are allowed and seem to agree with the data better than other alternatives. If these luminosity functions are still valid beyond our current detection limit, this would imply a sudden emergence of a large number of low-luminosity galaxies when looking back in time to $z\\approx 10$, which, while seemingly exotic, would naturally fit in the picture of the cosmic hydrogen reionization. These early galaxies could easily account for the ionizing photon budget required by the reionization, and they would imply that the global star formation rate density might start from a very high value at $z\\approx 10$, rapidly reach the minimum at $z\\approx 7$, and start to rise again towards $z\\approx 6$. In this scenario, the majority of the stellar mass that the universe assembled through the reionization epoch seems still undetected by current observations at $z\\approx 6$. ",
        "introduction": "An important subject in cosmology is the hydrogen reionization of the universe. The detections of complete Gunn-Peterson troughs (Gunn \\& Peterson 1965) in the spectra of a few $z>6$ QSOs (see Fan et al. 2006) from the Sloan Digital Sky Survey indicate that the reionization must have ended at around $z\\approx 6$, while the recent measurement of Thomson scattering optical depth from the seven-year Wilkinson Microwave Anisotropy Probe (WMAP) data shows that the reionization most likely began at $z=10.4\\pm1.2$ if assuming an instantaneous reionization history (Komatsu et al. 2010). An important and yet controversial question is the sources of reionization. It is clear that the QSO population can only produce a small fraction of the necessary ionizing photons at $z\\approx 6$ (e.g., Fan et al. 2002), which leaves star-forming galaxies as the most obvious alternative, although the measurement of the star-formation rate density (and hence the ionizing photon budget) is still uncertain (e.g., Stanway et al. 2003; Bouwens et al. 2003; Giavalisco et al. 2004b; Bunker et al. 2004; Yan \\& Windhorst 2004b). Yan \\& Windhorst (2004a; hereafter YW04a) pointed out that star-forming galaxies could account for the entire ionizing photon budget at $z\\approx 6$ as long as their luminosity function (LF) has a steep faint-end slope, $\\alpha<-1.6$. Using the Hubble Ultra Deep Field (HUDF; Beckwith 2006) data obtained by the Advanced Camera for Surveys (ACS), Yan \\& Windhorst (2004b; hereafter YW04b) have found a large number of $i_{775}$-dropouts, which are candidate galaxies at $z\\approx 6$, and obtained their LF that indeed has $\\alpha\\lesssim -1.8$. Such a very steep slope was later confirmed by Bouwens et al. (2006) using a larger $i_{775}$-dropout sample collected in the HUDF, the HUDF parallel fields and the two fields of the Great Observatories Origins Deep Survey (GOODS; Giavalisco et al. 2004a). Bouwens et al. (2007; hereafter B07) further compared the LF at $z\\approx 4$, 5 and 6, and suggested that LF has a strong luminosity evolution over this period in that $M^*$ is $\\sim 0.7$~mag fainter at $z\\approx 6$ than at $z\\approx 4$. On the other hand, they found that the LF does not change much in $\\alpha$.  As star-forming galaxies seem to be capable of keeping the universe ionized at $z\\approx 6$, it is natural to expect that star-forming galaxies are also capable of producing sufficient ionizing photons at earlier epochs to make the reionization happen, and therefore we should expect that they must exist in significant number extending well into the reionization epoch.  Along a different line of study, it has also been concluded that the universe must have started actively forming stars long before $z\\approx 6$. Using the GOODS {\\it Spitzer} Infrared Array Camera (IRAC) data in the HUDF region, Yan et al. (2005) have detected the restframe optical fluxes from three $z\\approx 6$ and eleven $z\\approx 5$ galaxies in the 3.6 and 4.5~$\\mu$m channels, and found that these are rather matured systems with stellar masses to the order of $\\sim 10^{10}M_\\odot$ and ages to the order of a few hundred Myr (see also Eyles et al. 2005). This strongly suggests that such high-mass galaxies must have started forming their stars at $z>7$ and likely earlier. Yan et al. (2006) further investigate this problem using a much larger sample in the entire GOODS field, and reinforced this conclusion (see also Eyles et al. 2007; Stark et al. 2007). Therefore, one should indeed expect a significant number of galaxies at $z\\gtrsim 7$.  To search for galaxies at $z\\gtrsim 7$, we have to carry out surveys in the near-IR regime, because the line-of-sight neutral hydrogen absorption effectively extincts the light from such sources that is emitted below 1~$\\mu m$ in observer's frame. In fact, we rely on this effect to identify such galaxies. The first candidate of this kind was reported by Dickinson et al. (2000) using the data obtained by the Near Infrared Camera and Multi-Object Spectrometer (NICMOS) No. 3 (NIC3) in the Hubble Deep Field, although now we have convincing evidence that it is likely a red galaxy at lower redshift, similar to those ``IRAC-selected Extremely Red Objects'' (IERO; Yan et al. 2004). Using the two broad-band data taken by NIC3 in the HUDF (Thompson et al. 2005), YW04b identified three objects that are missing from the ACS images. Subsequent analysis using the Spitzer IRAC data suggest that two of them are red galaxies at $z\\approx 2$--3 without much on-going star formation (Yan et al. 2004; but also see Mobasher et al. 2005 and Chary et al. 2006), while the third one (ID No.3) is less conclusive because it is blended with other sources in the IRAC image and thus its photometry is uncertain. Bouwens et al. (2004) used the same data set, and pushed to a fainter limit and found a few additional candidates. Bouwens et al. (2005) further included all available deep NIC3 imaging data and extended their search to $z\\approx 10$, although no conclusive answer was obtained. Using these results, Bouwens et al. (2008; hereafter B08) derived constraints to the UV LF of galaxies at $z\\approx 7$--10, and argued that the LF evolves strongly and continues the dimming trend in $L^*$ (as proposed in B06) to higher redshifts. As a result, the number density of galaxies, and hence the UV luminosity density in the earlier universe should be significantly lower than at a later time.  Meanwhile, the search for gravitationally lensed galaxies around foreground galaxy clusters have resulted in some remarkable success. Kneib et al. (2004) first discovered such an object that is likely at $z>6$. While no precise redshift was obtained, extensive optical and IR observations suggest that it is at $z\\approx 7$. Bradley et al.  (2008) reported a very bright, highly probable candidate at $z\\approx 7.6$. Zheng et al. (2009) reported three new $z\\approx 7$ candidates from the same campaign. Richard et al. (2008) have found 12 candidates at similar redshifts. These results, however, still do not allow us put a strong constraint to the number density of galaxies at these redshifts, although they indeed prove that very high-redshift galaxies much fainter than our current detection limits (should there be no lensing magnification) do exist. There are also some other evidence from the search for \\Lya emitters around foreground clusters that supports similar conclusion (Stark et al. 2007).  While the above results start to give us some meaningful constraints to the faint-end of the LF at $z\\gtrsim 7$, there is still only limited constraint to the bright-end where deep wide-field surveys are required. To date, most such surveys have null detection (e.g., Willis et al. 2008; Stanway et al. 2008; Sobral et al. 2009) or are uncertain (e.g. Hickey et al. 2009). One exception that has produced positive detections is the Subaru CCD survey for LAEs at $z=7.0\\pm0.1$, which has resulted in the redshift record of $z=6.96$ (Iye et al. 2006). However, even such LAEs can only constrain the faint population, as the vast majority of them are pure emission-line objects and are not seen in continuum. Another exception is that of Ouchi et al. (2009), which surveyed $\\sim 1500$~arcmin$^2$ in two fields and resulted in 22 $z\\approx 7$ candidates to $AB\\sim 26.0$~mag. Recently, Capak et al. (2009) reported three very bright candidates at $z\\approx 7$ at $J\\sim 23$~mag, however, their nature is still uncertain.  Nearly all these studies (with the exception of Capak et al. 2009) in the very high-redshift frontier have claimed that their results are consistent with the strong declining evolution of the LF with respect to increasing redshift as suggested by B07 and B08. If this is indeed what the universe behaves, we are facing a dilemma: on one hand, the hydrogen reionization, which is now well constrained that must have started at $z\\approx 10$, requires a large amount of strong UV emitting sources, and the significant stellar mass density measured at $z\\approx 6$ also strongly suggests very active star formation activities at $z\\gtrsim 8$ and above; on the other hand, the limited number of observations to search for galaxies beyond $z\\approx 7$ indicates a strong declining number density of galaxies at higher redshifts. To reconcile these seemingly conflict results, a more decisive survey for galaxies at $z\\gtrsim 7$ is in demand.  The Wide Field Camera 3 (WFC3) recently installed to \\hst\\, has provided a unique opportunity for the study of the universe at very high redshifts. The IR channel of this camera has a factor of $6.4\\times$ larger field-of-view (FOV) and an order of magnitude higher Q.E. as compared to NIC3, making it the most powerful tool in detecting galaxies at $z\\approx 7$ and beyond. For this reason, the first set of deep data that it took has already inspired four papers to appear at the arXiv preprint service within one week after the data were made public (Bouwens et al. 2009; Oesch et al. 2009; Bunker et al. 2009; McLure et al. 2009). All these new results, however, seem to reiterate that the number density of galaxies rapidly declines when we look back further in time. Here we present our results based on an independent reduction and analysis of these data. Our effort has resulted in more candidate galaxies at $z\\gtrsim 7$ than others, and, for the first time, a large sample of highly probable candidate galaxies at $z\\approx 10$. We will show that our analysis has led to a completely new, although still tentative conclusion about the formation and evolution of galaxies in the early universe.  Our paper is organized as following. We briefly describe the WFC3 IR instrument and its observations of the HUDF in \\S 2, and give the details of our data reduction in \\S 3. Our photometry and catalog construction is described in \\S 4. The candidate selection and the dropout samples are presented in \\S 5. We discuss the implications of our results in \\S 6, followed by a summary in \\S 7. For simplicity, we denote the ACS passbands F435W, F606W, F775W, and F850LP as $B_{435}$, $V_{606}$, $i_{775}$, and $z_{850}$, respectively, and denote the three WFC3 IR passbands F105W, F125W and F160W as $Y_{105}$, $J_{125}$ and $H_{160}$, respectively. All magnitudes quoted are in AB system. Throughout the paper, we use the following cosmological parameters: $\\Omega_M=0.27$, $\\Omega_\\Lambda=0.73$ and $H_0=71$~km~s$^{-1}$~Mpc$^{-1}$. ",
        "conclusions": ""
    },
    "0910/0910.3074_arXiv.txt": {
        "abstract": "We report the discovery of 31.18\\un{ms} pulsations from the \\integ~source \\igr~using the Rossi X-ray Timing Explorer (\\xte). This pulsar is most likely associated with the bright Chandra \\xray~point source lying at the center of \\snr, a previously unrecognised Galactic composite supernova remnant with a bright central non-thermal radio and \\xray~nebula, taken to be the pulsar wind nebula (PWN). \\psr~is amongst the most energetic rotation-powered pulsars in the Galaxy, with a spin-down luminosity of $\\dot E = 5.1 \\times 10^{37}$ erg~s$^{-1}$. In the rotating dipole model, the surface dipole magnetic field strength is $B_s = 1.1 \\times10^{12}$~G and the characteristic age $\\tau_c \\equiv P/2\\dot P = 12.7$~kyr. The high spin-down power is consistent with the hard spectral indices of the pulsar and the nebula of $1.22\\pm0.15$ and $1.83\\pm0.08$, respectively, and a 2--10\\un{keV} flux ratio $F_{PWN}/F_{PSR} \\sim 8$. Follow-up Parkes observations resulted in the detection of radio emission at 10 and 20\\un{cm} from \\psr~at a dispersion measure of $\\sim$ 560 cm$^{-3}$ pc, which implies a relatively large distance of 10 $\\pm$ 3\\un{kpc}. However, the resulting location off the Galactic Plane of $\\sim$ 280\\un{pc} would be much larger than the typical thickness of the molecular disk, and we argue that \\snr~lies at a distance of $\\sim$ 7\\un{kpc}. There is no gamma-ray counterpart to the nebula or pulsar in the Fermi data published so far. A multi-wavelength study of this new composite supernova remnant, from radio to very-high energy gamma-rays, suggests a young ($\\lesssim$ 10$^{3}$ yr) system, formed by a sub-energetic ($\\lesssim 10^{50}$ ergs), low ejecta mass (M$_{\\rm ej} \\sim 3$ \\msun) SN explosion that occurred in a low-density environment ($n_0 \\sim$ 0.01 cm$^{-3}$). ",
        "introduction": "\\label{s:intro}  The number of known supernova remnants (SNRs) in the Galaxy has increased significantly over the last few years, mainly due to a new generation of radio and \\xray~instruments of unprecedented sensitivities and angular resolutions. In particular, these observations, combined with Galactic Plane surveys \\cite[\\eg][]{c:brogan06}, and targeted observations \\cite[\\eg][]{c:gelfand07,c:gaensler08} have increased the fraction of SNRs found to harbor an energetic pulsar powering a wind nebula (the so-called composite SNRs). Moreover, deep observations toward compact PWNe have proven to be successful in detecting their powering pulsars \\citep{c:camilo09,c:gotthelf09}, providing constraints on their energetics, spin evolution, and birth parameters. Along with the current generation of high/very-high energy (VHE; $>$ 100\\un{GeV}) instruments, these observations are of prime importance for understanding the structure and evolution of these sources, and the underlying acceleration mechanisms which occur close to the pulsar, and at the relativistic and non-relativistic shock fronts bounding the PWN and the host SNR \\citep{c:gs06}.  The soft $\\gamma$-ray source \\igr~was discovered in a deep \\integ/\\ibisg~mosaic of the Circinus region as a persistent source at the mCrab level \\citep{c:keek06}. A \\swi/XRT survey of \\integ~sources located \\igr~to 4\\s~in 2--10\\un{keV} \\xrays, but no conclusion was reached on its nature \\citep{c:malizia07}. A follow-up \\chandra~survey of unidentified IGR sources reported a PWN within a $\\sim$ 3\\m~diameter nearly circular emission nebula \\citep{c:tomsick09}. These authors presented the 0.3--10\\un{keV} spectrum of the total emission by an absorbed power-law with a relatively hard photon index ($\\Gamma$ = 1.82 $\\pm$ 0.13) and a large column density \\nh~$\\sim$ 3 $\\times$ 10$^{22}$ cm$^{-2}$.  In this paper, we report the discovery with the Rossi X-ray Timing Explorer (\\xte) of 31.18\\un{ms} pulsations toward \\igr, and the radio detection in follow-up observations using the Parkes telescope. We present a multi-wavelength study of this new Galactic composite SNR, \\snr, using radio, X-ray, and gamma-ray data. \\psr~is one of the most energetic in the Galaxy and powers a wind nebula whose broadband non-thermal synchrotron emission is measured in radio and \\xrays. We present a spatial and spectral analysis that shows evidence for a young SNR, formed from a sub-energetic, low ejecta mass SN explosion that occurred in a low-density environment. We also discuss the implications of the lack of high-energy (HE) gamma-ray emission in the \\fermi~data. ",
        "conclusions": "Even though some assumptions have been made in the previous estimates (\\eg~constant PWN magnetic field and PSR spin-down power), they give valuable insight on the nature of this new Galactic composite SNR. \\snr~harbors a highly energetic 31.18\\un{ms} pulsar, \\psr, which powers a wind nebula, both lying at the center of the host shell. All of the existing multi-wavelength observations suggest it is a young SNR ($\\lesssim$ 10$^{3}$ yr), and most likely distant ($>$ 5\\un{kpc}). However, many questions still remain to be answered. First, the distance is not very well constrained, though a large distance is favored by the dispersion measure of the radio pulse emission and by the large \\nh~measured with \\chandra. High-resolution observations of the ISM tracers such as HI and $^{12}$CO are then warranted to assess the surrounding medium properties. Moreover, the break measured at 6\\un{keV} and the spectral softening at increasing distances in the PWN need to be confirmed, and the SNR \\xray~spectrum needs to be investigated with more \\xray~data. Nevertheless, \\snr~falls into the emerging class of ``multi-wavelength'', young ($\\tau \\lesssim$ a few 10$^{3}$ yr) and composite SNRs, harboring very energetic PSRs and wind nebulae shining in radio, \\xrays~and potentially in HE/VHE \\gammarays. The list includes Kes~75 and G21.5$-$0.9 \\citep{c:gotthelf00,c:camilo06,c:bietenholz08,c:djannati08}, and more recently G0.9+0.1 \\citep{c:aharonian05a} and HESS~J1813$-$178 \\citep{c:aharonian05b}, whose long-expected PSRs have recently been discovered \\citep{c:camilo09,c:gotthelf09}. It is of interest to note that \\fermi/LAT has not detected a bright source \\citep{c:fermi_soucat}, nor pulsar \\citep{c:fermi_psr1,c:fermi_psr2} coincident with \\snr, although, at first glance, energetic PSRs should be the most easily detectable sources. However, not all of the above-mentioned young and energetic PSRs have been detected by \\fermi. From the first \\fermi~catalog of gamma-ray PSRs, the sensitivity for a blind search of pulsed emission in the Galactic Plane is conservatively taken to be 2 $\\times$ 10$^{-7}$ cm$^{-2}$ s$^{-1}$ above 100\\un{MeV} \\citep{c:fermi_psr2}. Assuming a spectrum similar to that of PSR~J1833$-$1034 associated with G21.5$-$0.9, \\psr~features a maximal efficiency $\\eta =\\,L_{\\gamma}$/\\edot~of $\\sim$ 1 $d_{7}^2$ \\%. This is close to what is measured from other young PSRs (see Fig. 6 of \\citet{c:fermi_psr2}), and argues in favor of a fairly large distance to the source, as outlined in section \\ref{s:discu}. Further HE/VHE observations of \\snr~will certainly provide important constraints both on the PSR gamma-ray spectrum and on the PWN magnetic field strength."
    },
    "0910/0910.1592_arXiv.txt": {
        "abstract": "Here we show that the overabundance of ultra-luminous, compact X-ray sources (ULXs) associated with moderately young clusters in interacting galaxies such as the Antennae and Cartwheel can be given an alternative explanation that does not involve the presence of intermediate mass black holes (IMBHs). We argue that gas density within these systems is enhanced by the collective potential of the cluster prior to being accreted onto the individual cluster members and, as a result, the aggregate X-ray luminosity arising from the neutron star cluster members can exceed $>10^{39}\\;{\\rm erg s^{-1}}$. Various observational tests to distinguish between IMBHs and accreting neutron star cusps are discussed. ",
        "introduction": "Over the years, the existence of two distinct populations of black holes has been established beyond a reasonable doubt.  Supermassive black holes, $M > 10^6 \\, M_\\sun$, are inferred in many galactic centers \\citep{1995ARAA..33..581K,1998AJ....115.2285M}, while stellar mass black holes, $M \\sim 1-10 \\, M_\\sun$, have been identified by their interaction with companion stars \\citep{2006csxs.book..157M}. The situation at intermediate masses, $M \\sim 10^2 - 10^5 M_\\sun$, is still uncertain despite recent evidence for mass concentrations within the central regions of some globular clusters \\citep{2005ApJ...634.1093G,2007ApJ...661L.151U,2008ApJ...676.1008N}. This evidence remains controversial, partly because the velocity dispersion profiles can be reproduced without invoking the presence of an intermediate mass black hole (IMBH) \\citep{baum1, baum2,2009arXiv0905.0627A}.  Recently, some evidence has arisen for the presence of IMBHs in moderately young star clusters, where ultra-luminous, compact X-ray sources (ULXs) have been preferentially found to occur \\citep{2001ApJ...554.1035F,2008AIPC.1010..357T}.  Their high luminosities have been interpreted as imprints of IMBHs \\citep{2004Natur.428..724P}, rather than binaries containing a normal stellar mass black hole \\citep{2006ApJS..166..211Z}. In this {\\it Letter}, we present an alternative explanation for the overabundance of ULXs associated with young clusters in galaxies such as the Antennae and Cartwheel.  In this new paradigm, the accretion of gas by the collective star cluster potential moving through the merging medium is strongly enhanced relative to the individual rates and, as a result, the aggregate X-ray luminosity arising from the neutron star cluster members can exceed $>10^{39}\\;{\\rm erg s^{-1}}$.  Much of the effort herein will be dedicated to understanding the conditions by which the collective potential of a star cluster is able to accrete gas with highly enhanced rates and its effect on the integrated accretion luminosity of the neutron star cluster members.  The conditions found in systems such as the Antennae galaxy, as we will argue, are favorable for this type of mechanism to operate effectively and produce an overabundance of ULXs. ",
        "conclusions": "Many studies have been focused on the flow around compact stars with a point mass potential. Although clusters have much larger masses than individual stars, their potential is relatively shallow. In this paper we consider the efficiency of accretion in these cluster potentials, and show that when the sound speed or the relative velocity of the ambient medium is less than the central velocity dispersion of the cluster, the collective potential alters the local gas flow before the gas is accreted onto the individual stars within the cluster. Accretion onto these dense stellar cores at the inferred rate can lead to the onset of ULX sources.  While there are no stellar clusters observed in the Galactic disk which bear these anticipated properties (the relative velocity of the halo clusters to the interstellar medium is in the range of 100 km/s), observations of several cluster knots in the Antennae indicate intracluster relative velocities that comparable to the central velocity dispersions \\citep{2005AJ....130.2104W}.  Based on the results of the current work, we show that accretion by individual compact stars in the centers of such systems is enhanced greatly relative to their rate of accretion directly from the ambient gas, and conclude that this process may be relevant for explaining the origin of ULX sources in these extraordinary clusters.  Illumination of nearby gas clouds by these sources may also lead to reprocessed infrared, optical and ultraviolet emission.  Finally, the sources may leave trails of denser and likely hotter gas behind them as they plough through the gas.  A way to distinguish between an IMBH \\citep{2004Natur.428..724P} and an accreting neutron star cusp is via time-dependent observations. The emission from a relativistic region of a IMBH might vary on time-scales of seconds.  The emission of a large number of statistically independent black holes and neutron stars should be considerably less variable: $\\Delta t \\leq r_{\\rm c}/c_{\\rm s}\\sim 10^4$ yr. Observations find that a handful of the X-ray sources in the Antennae galaxy are indeed variable albeit on timescales that are larger than a few years \\citep{2006ApJS..166..211Z}. Such variability might be explained if only a moderate fraction of the compact stars dominate the total luminosity. Such compact isolated accretors will probably have unusual time-variability properties as their discs may be much larger than the typical discs of X-ray binaries, and indeed they are missing the perturbing influence of the secondary. On the other hand, accretion disc feeding in these sources will be variable itself, leading to variability on variety of time-scales.  Although accretion disk spectra are notoriously difficult to calculate from first principles, an IMBH and a cluster core may also have observably different spectra. It has been suggested that a cool multicolor disk spectral component, might indicate the presence of an IMBH \\citep{2003ATel..212....1M}.  This is understood as following. The larger mass of the BH accretor, the lower the temperature of the inner edge of the disk, which scales as $T_{\\rm i} \\propto (M_{\\rm BH}\\dot{M}/r_{\\rm i})^{1/4} \\propto \\dot{M}^{1/4}M_{\\rm BH}^{-1/2}$ for a simple thin-disk model. Since for our model, individual neutron star sources have $T_{\\rm i}\\sim 1$ keV (consistent with observations), then an IMBH might have $T_i \\sim 30$ eV. Thus, an association of the ULXs with an IMBH, as opposed to a accreting distribution of compact remnants, could be made on the basis of a very soft observed spectral component.  We know of no reported observations of such components."
    },
    "0910/0910.1906_arXiv.txt": {
        "abstract": "Dark energy must cluster in order to be consistent with the equivalence principle. The background evolution can be effectively modeled by either a scalar field or by a barotropic fluid. The fluid model can be used to emulate perturbations in a scalar field model of dark energy, though this model breaks down at large scales. In this paper we study evolution of dark energy perturbations in canonical scalar field models: the classes of thawing and freezing models. The dark energy equation of state evolves differently in these classes. In freezing models, the equation of state deviates from that of a cosmological constant at early times. For thawing models, the dark energy equation of state remains near that of the cosmological constant at early times and begins to deviate from it only at late times. Since the dark energy equation of state evolves differently in these classes, the dark energy perturbations too evolve differently. In freezing models, since the equation of state deviates from that of a cosmological constant at early times, there is a significant difference in evolution of matter perturbations from those in the cosmological constant model. In comparison, matter perturbations in thawing models differ from the cosmological constant only at late times. This difference provides an additional handle to distinguish between these classes of models and this difference should manifest itself in the ISW effect. ",
        "introduction": "Various observation have confirmed that the expansion rate of the universe is accelerating \\cite{obs_proof}. These observations include those of Supernova type Ia \\cite{nova_data}, observations of Cosmic Microwave Background \\cite{boomerang,wmap_params} and large scale structure \\cite{sdss}. The accelerated expansion of the universe can  be explained by introducing a cosmological constant $\\Lambda$ in the Einstein's equation \\cite{ccprob_wein,review3}. However, the cosmological constant model is plagued by the fine tuning problem \\cite{ccprob_wein}. This has motivated the study of dark energy models to explain the current accelerated expansion of the universe (for reviews see \\cite{DEreview}). An alternative to the cosmological constant model is to assume that this accelerated expansion is driven by a canonical scalar field with a potential $V(\\phi)$, namely the quintessence field \\cite{quint1,expo,linear,quadratic,invphi,invexpo}. There exists another class of string theory inspired scalar field dark energy models known as tachyon models \\cite{tachyon1,2003PhRvD..67f3504B} and there are models which allow $w<-1$ are known as phantom models \\cite{2002PhLB..545...23C}. Phantom type dark energy can also be realized in a scalar tensor theory of gravitation \\cite{STG}. Other scalar field models include k-essence field \\cite{2001PhRvD..63j3510A}, branes \\cite{brane1} and fluid models like the Chaplygin gas model and its generalizations \\cite{chaply}. There are also some phenomenological models \\cite{water}, field theoretical and renormalization group based models  (see e.g. \\cite{tp173}), models that unify dark matter and dark energy \\cite{unified_dedm1}, holographic dark energy models \\cite{HGDE}, QCD dark energy \\cite{qcddark}  and many others like those based on horizon thermodynamics (e.g. see \\cite{2005astro.ph..5133S}).     Different models of dark energy which have the same background evolution are indistinguishable purely from distance measurements. Evolution of perturbations in these models is expected to break this degeneracy. The Integrated Sachs Wolfe (ISW) effect can distinguish a cosmological constant from other models of dark energy, especially ones with a dynamical dark energy \\cite{ddw}. Dark energy perturbations have been extensively studied in the linear approximation \\cite{weller_lewis,bean_dore,depert,gordonhu}. Perturbations in dark energy affect the low $l$ quadrupole in the CMB angular power spectrum through the ISW effect \\cite{weller_lewis,bean_dore}. It was shown in Ref.\\cite{weller_lewis} that dark energy perturbations affect the low $l$ quadrupole in the CMB angular power spectrum through the ISW effect. For models with $w>-1$ this effect is enhanced while for phantom like models it is suppressed. In these models dark matter perturbations and dark energy perturbations are anti-correlated for large effective sound speeds. This anti-correlation is a gauge dependent effect \\cite{bean_dore}. There are several other studies of perturbations in dark energy \\cite{chpgas_pert}, including some that deal with evolution of spherical perturbations \\cite{sph_coll,mota}. For canonical scalar field dark energy, the perturbations in matter are enhanced by the presence of dark energy perturbations in comparison with smooth dark energy model \\cite{ujs}. The matter perturbations in fluid models are suppressed compared to corresponding homogeneous dark energy scenario \\cite{hkj}. As long as the speed of propagation of perturbations `$c_s^2$' is positive the evolution of matter perturbations is indistinguishable from a smooth dark energy model. This is true for scales smaller then the Hubble radius. Dark energy perturbations in a fluid model with an appropriate $c_s^2$ emulate that of a scalar field model very well below the Hubble scale but start to differ at larger scales. Therefore the fluid model is not a good approximation at these scales. This also implies that the growth of perturbations at large scales depends on the details of the model even though the background evolution is the same. A separate analysis is therefore required for every  model.    In this paper we consider different scalar field models to study evolution of dark energy perturbations. We consider two different types of potentials, classified as '{\\it thawing}' and '{\\it freezing}' in Ref. \\cite{limitsofq}. For potentials with a thawing behavior, the scalar field is frozen at early times and starts to roll down the potential at late time. Hence the equation of state of dark energy starts near $w=-1$ at early times to $w>-1$ at late times. In contrast, in the case of potentials with freezing behavior, the scalar field rolls down the potential and approaches the minimum of the potential, with the equation of state going from $w>-1$ to freezing at $w=-1$ at late times. Due to the different evolution of the equation of state, it is expected that the perturbations in dark energy evolve differently in these classes. For freezing potentials, at early times, the equation of state deviates from $w=-1$ hence the perturbations in dark energy are expected to be enhanced. Whereas, in thawing potential scenarios, dark energy perturbations become significant only at late times. Since dark energy perturbations enhance perturbations in nonrelativistic matter, models with  freezing type potentials will have more enhanced perturbations in matter  than those with thawing type behavior. This can be an additional tool (apart from distance measurements) to distinguish between these types of models.   The paper is organized as follows. In Section \\ref{sec::fg} we discuss evolution of perturbations in scalar field potentials in the thawing and freezing classes. The results are summarized in the concluding Section \\ref{sec::concl}. ",
        "conclusions": "\\label{sec::concl}  In this paper, we analyze the growth of perturbations in scalar field dark energy scenarios. The assumption that the distribution of dark energy (with $w \\neq -1$) is homogeneous at all length scales is inconsistent with the equivalence principle. On length scales comparable to or greater than the Hubble radius, the perturbations in dark energy can become comparable to perturbation in matter if $w_{de} \\neq -1$. For scales smaller than the Hubble radius, perturbations in dark energy can be neglected in comparison with the perturbation in matter at least in the  linear regime. Hence any deviations from the cosmological constant model can be explored by assuming dark energy as a homogeneous component and the scalar field models can be approximated well by parameterized fluid models. However, on much larger scales, if the equation of state parameter deviates from -1, then perturbations in dark energy do influence  matter power spectrum.   For canonical scalar fields, a clustering dark energy  enhances matter perturbations as compared to the corresponding homogeneous dark energy scenario. This enhancement more pronounced at scales larger than the Hubble radius. In particular, we study two broad classes of canonical scalar field models, namely the thawing and freezing models. The equation of state of dark energy evolves differently in these two classes of models. This affects the growth of perturbations in these different types. For fast rolling, freezing type models, the dark energy equation of state deviates away from that of a cosmological constant at early times and freezes to $w=-1$ at late times. In slow roll thawing models, the scalar field remains at a constant $w=-1$ and starts to  deviate away from this value at late times. The present day observations, which are based on distance measurements, cannot distinguish between these models.   We studied evolution of matter perturbations in the presence of a clustering dark energy. The models we have considered in this paper are within the range allowed by distance measurement observations. For  corresponding homogeneous dark energy models and clustering dark energy models, there are significant changes in the matter density contrast evolution. All the canonical scalar field models studied here show an enhancement in matter perturbations if dark energy is perturbed. Although we have not considered phantom like models in this paper, it is worth mentioning that in these models, dark energy perturbations suppress matter perturbations.   In general, freezing type models have a higher rate of growth of density contrast at early times. This is due to the fact the in these models, the equation of state of dark energy is further away from that of a cosmological constant. Hence dark energy perturbations play a prominent role at earlier times. For thawing type models, in general, dark energy perturbations affect the matter perturbations at late times. For most of the evolution, the matter perturbations remain close to those in cosmological constant model and being to deviate as the field begins to thaw. There are significant deviations in the way density contrast grows, not only between various  models but also  from the concordant cosmological constant model. Apart from  different way matter perturbations grow in the freezing and thawing classes of models, models within the same class also have variation in the behavior of the density contrast. Observable changes in the angular power spectrum at large scales are limited by the cosmic variance and therefore CMB data alone will be insufficient to distinguish between these models. Since the scales at which dark energy perturbations are relevant are large, the primary contribution to CMB anisotropies is through the ISW effect. Therefore it is important to study the ISW cross correlation with large scale structure indicators."
    },
    "0910/0910.4238_arXiv.txt": {
        "abstract": " ",
        "introduction": "The role of the gas in the formation and evolution process of early-type galaxies is still not fully understood. For example, recent studies have shown a large complexity in the gas structures in these systems (e.g. Morganti et al. 2006, Sarzi et al. 2006, Combes et al 2007) despite their often unspectacular optical appearance. In addition, some sources show nuclear activity while others do not. Cold-gas structures represent a fossil record of the formation and evolution of early-type galaxies. In particular, gas found on kiloparsec scales can be used to trace the evolution of the host galaxy (e.g. major merger vs. small accretions). In that respect, neutral hydrogen is an important tracer of these events as it often extends the dust and optical (disk) structure by a factor of two or more.  Close to the centre of galaxies ($<100$~pc), (cold) gas also plays a crucial role as it can provide the fuel that is needed to make the central black hole active. It is often believed that mergers are important in driving gas to the centre (see e.g. Hibbard \\& van Gorkom 1996, Barnes 2002), but recent studies of radio galaxies have shown that the activity in some galaxies may be associated with the (slow) accretion of gas from the ISM/IGM (e.g. Best et al. 2005).  Despite the need for large statistical investigations to understand (and classify) the different mechanisms at work (e.g. mergers, interactions, accretion, AGN activity etc.), it is indispensable to study close-by objects in great detail with very high linear resolution. For example, the circumnuclear region around the black hole ($< 100$~pc) can only be resolved in the most nearby galaxies to a degree that is needed to understand the accretion/fueling process.  By far the closest radio-loud early-type galaxy is Centaurus~A (NGC~5128) at a distance of only 3.8~Mpc\\footnote{At this distance 1\\arcmin ~corresponds to $\\sim 1.1$~kpc.} (Harris 2009). Cen~A has been studied in all possible wavelength regimes with very high linear resolution (for an overview see the review by Israel (1998) and the other contributions to this volume). However, it is crucial to compare Cen~A with other sources to check whether it is a typical example of its class, or whether it is unusual w.r.t. other radio-loud, low-luminosity sources.  In this paper we discuss whether Cen~A is special as seen from the neutral hydrogen perspective. That is, do the Cen~A properties (such as e.g. \\HI mass, morphology, kinematics etc.) differ from other early-type galaxies and does Cen~A share properties with radio galaxies that have comparable luminosity? In Sect.~2 we give a brief description of the \\HI morphology and kinematics on kpc and sub-kpc scales. Section~3 compares Cen~A with other nearby early-type galaxies and in Sect.~4 Cen~A is compared to a complete sample of nearby radio galaxies. We summarize our comparison in Sect.~5. ",
        "conclusions": "In order to understand Centaurus~A in the context of galaxy formation and evolution, we have compared the \\HI properties of Cen~A with early-type and radio galaxies. The \\HI mass, its distribution and the mainly settled kinematics is commonly found in other early-type/radio galaxies. The current phase of AGN activity is not connected to the recent merger which is in line with recent results for a sample of radio galaxies. The absorption against the nucleus is red- and blueshifted with respect to the systemic velocity and is in agreement --- as is also the case in other sources --- with a circumnuclear \\HI disk/torus structure. Hence, Centaurus~A seems to be --- from an \\HI perspective --- a typical galaxy of its class."
    },
    "0910/0910.2714_arXiv.txt": {
        "abstract": "{Current stellar population models have arguably the largest uncertainties in the near-IR wavelength range, partly due to a lack of large and well calibrated empirical spectral libraries. In this paper we present a project, which aim it is to provide the first library of luminosity weighted integrated near-IR spectra of globular clusters to be used to test the current stellar population models and serve as calibrators for the future ones. Our pilot study presents spatially integrated $K$-band spectra of three old ($\\ge$10~Gyr) and metal poor ([Fe/H]$\\sim$~--1.4), and three intermediate age (1~--~2~Gyr) and more metal rich ([Fe/H]$\\sim$~--0.4) globular clusters in the LMC. We measured the line strengths of the \\na\\/, \\ca\\/ and \\co\\/ absorption features. The \\na\\/ index decreases with the increasing age and decreasing metallicity of the clusters. The \\dco\\/ index, used to measure the \\co\\/ line strength, is significantly reduced by the presence of carbon-rich TP-AGB stars in the globular clusters with age $\\sim$1~Gyr. This is in contradiction with the predictions of the stellar population models of Maraston (2005). We find that this disagreement is due to the different CO absorption strength of carbon-rich Milky Way TP-AGB stars used in the models and the LMC carbon stars in our sample. For globular clusters with age $\\geq$2~Gyr we find \\dco\\/ index measurements consistent with the model predictions.} ",
        "introduction": "\\label{sec:motivation}  Since the 90's, the interpretation of the integrated light of galaxies (in the nearby universe or at high redshift) heavily relies on evolutionary population synthesis (EPS) models. Such models were pioneered by \\citet{tinsley80} and the method has been seriously extended since then \\citep[e.g.][]{bc93,worthey94,vazdekis96,fioc97,starburst99, maraston05, schiavon07}. They are used to determine ages, element abundances, stellar masses, stellar mass functions, etc., of those stellar populations that are not resolvable into single stars with today's instrumentation, i.e. most of the universe outside the Local Group. To build such EPS models we use simple stellar populations (SSP). There are two essential advantages of focusing on SSPs. First, SSPs can be reliably calibrated. They can be compared directly with nearby globular cluster (GC) data for which accurate ages and element abundances are independently known from studies of the resolved stars. This step is crucial to fix the stellar population model parameters that are used to describe model input physics and which cannot be derived from first principles (e.g., convection, mass loss and mixing).  Second, SSPs can be used to build more complex stellar systems. Systems made up by various stellar generations can be modelled by convolving SSPs with the adopted star formation history \\citep[e.g.][] {kodama97, bc03}. Models describing accurately the integrated light properties, including medium to high resolution spectra and/or line-strength indices, are and will be our main tool to investigate and analyse the star-formation history over cosmological time-scales.  This approach has worked well in the optical spectroscopic regime and has led to well calibrated models \\citep[e.g.][]{tmb03,bc03,maraston05}. With the application of such models to observed spectra we derive reasonable estimates of the main stellar population parameters (age, chemical composition and M/L ratio) in the nearby universe \\citep[e.g.,][]{kun00,trager00,thomas05,cappellari06,sb07} as well as at higher redshifts \\citep[e.g.,][]{bernardi05,maraston05,sb09}. Of course, uncertainties remain due to the degeneracy of age and metallicity effects in the optical wavelength range \\citep[e.g.,][] {worthey94}.  The integrated near-IR light in stellar populations with ages $\\geq$ 1 Gyr is dominated by one stellar component, cool giant stars, whose colour and line indices are mainly driven by one parameter: metallicity \\citep{frogel78}.  Near-IR colours and indices also have the advantage of being more nearly mass-weighted, i.e. the near-IR mass-to-light ratio is closer to one \\citep [see e.g.,][]{worthey94}. So, by combining the optical as well as near-IR information one can resolve the currently remaining degeneracies between age and chemical composition, present in the models, and hope to gain a better understanding of star-formation histories. However, currently available stellar population models have arguably the largest uncertainties in the near-IR and thus it is of paramount importance to provide high-quality observational data to validate and improve the state-of-the-art models.  Globular clusters in the Local Group are an ideal laboratory for this project since ample information from studies of the resolved stars is available. Yet, integrated spectroscopic observations of the Galactic GCs in the near-IR are very challenging due to their large apparent sizes on the sky. The Large Magellanic Cloud (LMC) and its globular cluster system, located about 50 kpc away, is a much better observational choice. It shows evidence for a very complex and still ongoing star formation activity. The LMC GCs have an advantage (for the scope of this project) with respect to Galactic GCs - they span a larger range in ages. Studies of the LMC globular cluster system show one old component with age $>$10~Gyr. After this time there was a ''dark age'' with just one cluster formed before a new burst of cluster formation that has started around $3-4$~Gyr ago \\citep{dacosta91}. A disadvantage is their lower metallicity. \\begin{center} \\begin{table*}[tdp] \\centering \\caption{\\label{tab:lmc_observations}Target globular clusters in the LMC -- observing log.} \\begin{tabular}{c c c c c c c } \\hline \\hline Name & UT date &  $V$ & $(B-V)$  & SWB & Age & [Fe/H]\\\\ (1) & (2) & (3) & (4) & (5) & (6) & (7) \\\\ \\hline NGC\\,1754 & 2006 Nov 18 & 11.57 & 0.75 & VII & 10 & -1.42$^{a}$, -1.54$^{b}$\\\\ NGC\\,2005 & 2006 Nov 29 & 11.57 & 0.73 & VII & 10 & -1.35$^{a}$, -1.92$^{b}$, -1.80$^{c}$, -1.33$^{d}$\\\\ NGC\\,2019 & 2006 Dec 10 & 10.86 & 0.76 & VII & 10 & -1.23$^{a}$, -1.18$^{b}$, -1.37$^{c}$, -1.10$^{d}$\\\\ NGC\\,1806 & 2006 Dec 06 & 11.10 & 0.73 & V   & 1.1& -0.23$^{b}$, -0.71$^{e}$ \\\\ NGC\\,2162 & 2006 Dec 05 & 12.70 & 0.68 & V   &  1.1  & -0.23$^{b}$, -0.46$^{f}$\\\\ NGC\\,2173 & 2006 Dec 06 & 11.88 & 0.82 & V-VI & 2& -0.24$^{b}$, -0.42$^{f}$, -0.51$^{g}$\\\\ \\hline \\hline \\end{tabular}  \\smallskip \\flushleft Notes: (1) Cluster name, (2) date of observation, (3) integrated $V$-band magnitude, (4) $(B-V)$ colour and (5) SWB type taken from \\citep{bica96,bica99}, (6) Age of the cluster in Gyr, based on the SWB type \\citep{frogel90}, (7) [Fe/H] derived using different methods: $^{a}$ \\citet{olsen98} -- slope of the RGB; $^{b}$ \\citet{olszewski91} -- low-resolution Ca\\,II triplet; $^{c}$ \\citet{johnson06} -- high-resolution Fe\\,I; $^{d}$~\\citet{johnson06} -- high-resolution Fe\\,II; $^{e}$ \\citet{dirsch00} -- Str\\\"{o}mgren photometry; $^{f}$ \\citet{groch06} -- low-resolution Ca\\,II triplet; $^{g}$~\\citet{muc08} -- high-resolution spectroscopy. \\end{table*} \\end{center} The goal of this project is to provide an empirical near-IR library of spectra for integrated stellar populations with ages $\\geq$ 1 Gyr, which will be used to verify the predictions of current SSP models in the near-IR wavelength range. Here we present the results from a pilot study of $K$-band spectra of 6 globular clusters in the LMC. The analysis of their $J$ and $H$-band spectra will be discussed in a separate paper.  This paper is organised as follows: in Sect.~\\ref{sec:sample} we give details about the sample selection and the observing strategy. Sect.~\\ref{sec:obs} is devoted to the observations and data reduction. In Sect.~\\ref{sec:in_or_out} we discuss the cluster membership of the stars in our sample, in Sect.~\\ref{sec:indices} we describe the near-IR index measurement procedures. In Sect.~\\ref{sec:data_models_lmc} we make a comparison between the currently available stellar population models in the near-IR with our data, discuss the observed disagreements, and give potential explanations. Finally, in Sect.~\\ref{sec:conclusions} we give our concluding remarks. ",
        "conclusions": "\\label{sec:conclusions}  The goal of this project is to provide an empirical spectral library in the near-IR for integrated stellar populations with ages $>$~1~Gyr, which will be used to test the current and calibrate the future stellar population models.  In this paper we have presented the first results from a pilot study of the $K$-band spectroscopic properties of a sample of six globular clusters in the LMC. To validate the observational strategy, data reduction and analysis methods, we have selected from the catalogue of \\citet{bica99} three out of 38 GCs with SWB type VII to represent the old ($>$10~Gyr) and metal poor ([Fe/H]$\\sim$--1.4) population of the LMC, and three out of 71 clusters with SWB types V and VI to explore the properties of the intermediate age (1\\,--\\,3~Gyr) and more metal rich ([Fe/H]$\\sim$--0.4) component of the population. For each cluster our integrated spectroscopy covers the central $24\\arcsec\\times24\\arcsec$ and in most of the cases we have sampled about half the light. However, in order to better sample bright AGB stars, which are the most important contributors to the integrated cluster light in the near-IR, we have observed up to 9 of the brightest stars outside the central mosaics, but still within the tidal radii of the clusters, that have near-IR colours and magnitudes consistent with bright red giants in the observed clusters. We obtained integrated luminosity weighted spectra for the six clusters, measured the line strengths of \\na\\/, \\ca\\/ and \\co\\/ absorption features in the $K$-band and compared the strength of \\co\\/ with the stellar population models of \\citet{maraston05}.  The observing strategy to cover at least the central half-light radius with a number of SINFONI pointings showed to be an efficient way in sampling the near-IR light of old ($>$\\,10\\,Gyr) clusters. For the intermediate age and sparse clusters, which are dominated by just a few very bright stars, observing a central mosaic plus a number of the brightest stars in the vicinity of the cluster is a better choice for optimal cluster light sampling. The availability of high spectral resolution spectroscopy greatly helps the differentiation of the cluster member stars from the LMC field population.  In intermediate age clusters the largest amount of light originates from oxygen (M-type) and carbon-rich (C-type) AGB stars. Different ratios of the contributions by these two types of stars can lead to significant changes in the near-IR \\co\\/ line strength. According to our observations, when the C-type stars contribution peaks (at $\\sim$1\\,Gyr) the observed CO line strength is weak and then increases fast to reach its maximum for clusters with age $\\sim$2\\,Gyr. It is important to note that a weak line strength of \\co\\/ does not mean that there is less CO in these clusters/stars. The indices, which are used to describe the line strengths, take also into account the continuum shape, which in the case of carbon-rich stars is severely affected by typical for this type of stars absorption features and thus the resultant index value is low.  The comparison of our data with the stellar population models of \\citet{maraston05}, in terms of CO line strength, shows a disagreement for the youngest clusters in our sample. For clusters with age $\\sim$1\\,Gyr the models predict maximal CO line strength, while we observe the opposite: the CO strength is significantly weaker. At the same time, literature data of the integrated colours of the clusters are consistent with these models. We explain these discrepancies as due to the different origin of the C-type stars used to calibrate the models and the ones in our data sample. The stars, used for model calibration, are Milky Way carbon stars, while our carbon stars are born in the LMC. We support this scenario with Fig.~\\ref{fig:co_jk_cstars}, where we show that carbon rich stars in the Milky Way and LMC, which have similar $(J-K)$ colour, have very different CO line strengths.  This deserves further investigation and hopefully the next generation of carbon star models \\citep[e.g.][]{aringer09} will help us to elucidate whether a systematic effect of the metallicity on CO indices is expected, or the found  discrepancy is due to the small sample of individual observations of variable stars in existing libraries.  The near-IR \\na\\/ index shows a dependance on the age of the clusters -- it is decreasing with an increasing age. The combination between optical and near-IR spectral indices seems to offer possibilities to break the age-metallicity degeneracy, but more accurate and detailed stellar population models are necessary in the near-IR wavelength range.  These models are of paramount importance, when studying the spatially resolved stellar populations of nearby galaxies. Adaptive optics assisted observations allow for the best correction of the Earth's atmosphere perturbing effects when observing in the near-IR. First attempts to explore galaxy evolution via spatially resolved near-IR spectroscopy \\citep[e.g.][]{az08,davidge08,nowak08} have shown the need for a better understanding of the properties of stellar populations in this wavelength range. The availability of detailed and reliable stellar population models in the near-IR will open up a new window to the exploration of galaxy formation and evolution.  The final integrated spectra of the six LMC globular clusters will be made available via the Strasbourg astronomical Data Center (CDS)."
    },
    "0910/0910.0711_arXiv.txt": {
        "abstract": "We have used a statistical technique ``Shuffle\" \\citep {bhav, bharad2} in seven nearly two dimensional strips from the Sloan Digital Sky Survey Data Release Six (SDSS DR6) to test if the statistically significant length scale of filaments depends on luminosity, colour and morphology of galaxies. We find that although the average filamentarity depends on these galaxy properties, the statistically significant length scale of filaments does not depend on them. We compare it's measured values in SDSS against the predictions of $\\Lambda$CDM N-body simulations and find that $\\Lambda$CDM model is consistent with observations. The average filamentarity is known to be very sensitive to the bias parameter. Using $\\Lambda$CDM N-body simulations we simulate mock galaxy distributions for SDSS NGP equatorial strip for different biases and test if the statistically significant length scale of filaments depends on bias. We find that statistically significant length scale of filaments is nearly independent of bias. This result is possibly related to the fact that statistically significant length scale of filaments is nearly the same for different class of galaxies which are differently biased with respect to underlying dark matter distribution. The average filamentarity is also known to be dependent on the galaxy number density and size of the samples. We use $\\Lambda$CDM dark matter N-body simulations to test if the statistically significant length scale of filaments depends on number density of galaxies and size of the samples. Our analysis shows that the statistically significant length scale of filaments very weakly depends on these factors. Finally we test the reliability of our method by applying it to controlled samples of segment Cox process and find that our method successfully recovers the length of the inputted segments. Summarizing these results we conclude that the statistically significant length scale of filaments is a robust measure of the galaxy distribution. ",
        "introduction": "The fact that the galaxies appear to be distributed along filaments which are interconnected to form a web like structure which is often reffered as the `cosmic web' is one of the most striking visual feature in all the present and past redshift surveys (e.g. , CfA , \\citealt{gel}; LCRS, \\citealt{shect}; 2dFGRS, \\citealt{colles} and SDSS, \\citealt{stout}).  The analysis of filamentary patterns in the galaxy distribution has a long history dating back to a few papers in the late-seventies and mid-eighties by \\citet{joe}, \\citet{einas4}, \\citet{zel}, \\citet{shand1} and \\citet{einas1}. Filaments are the most striking visible patterns seen in the galaxy distribution (e.g. \\citealt{gel}, \\citealt{shect}, \\citealt{shand2}, \\citealt{bharad1}, \\citealt{mul}, \\citealt{basil}, \\citealt{doro2}, \\citealt{pimb}, \\citealt{pimb1}, \\citealt{pandey}). A review on a number of physically motivated and statistical methods to define filaments is provided in \\citet{pimb2}. The percolation analysis (eg. \\citealt{shand1}, \\citealt{einas1}), the genus statistics (eg. \\citealt{gott}), the minimal spanning tree (e.g. \\citealt{barrow}), the Voronoi tessellation (\\citealt{ike}, \\citealt{weygaert}), the Minkowski functionals (eg. \\citealt{mecke}, \\citealt{smal}) and the `Shapefinders' \\citep{sahni} are some of the useful statistical tools introduced to quantify the topology and geometry of the galaxy distribution. \\citet{stoi1} propose to apply a marked point process to automatically delineate filaments in the galaxy distribution. \\citet{colberg} studied the intercluster filamentary network in high resolution N-body simulations of structure formation in a $\\Lambda$CDM Universe. \\citet{arag} use Multiscale Morphology Filter technique to identify wall-like and filament-like structures in cosmological N-body simulations.  \\citet{sus} propose a skeleton formalism to quantify the filamentary structure in three dimensional density fields. \\citet{stoi2} propose to apply an object point process to objectively identify filaments in galaxy redshift surveys. \\citet{sarkar1} propose the Local Dimension to locally quantify the shape of large scale structures in the neighbourhood of different galaxies in the Cosmic Web.  The SDSS \\citep{york} is currently the largest galaxy redshift survey. In an earlier work \\citep{pandey} (hereafter Paper I) we have analysed the filamentarity in the equatorial strips of this survey. These strips are nearly two dimensional and we have projected the data onto a plane and analysed the resulting 2-D galaxy distribution. We find evidence for connectivity and filamentarity in excess of that of a random point distribution, indicating the existence of an interconnected network of filaments. We find that filaments are statistically significant upto length scales $80 \\, h^{-1} {\\rm Mpc}$ and not beyond \\citep{pandey}. All the structures spanning length-scales larger than this length scale are the result of chance alignments. This is consistent with an earlier analysis by \\citet{bharad2} where they show that in Las Campanas Redshift Survey (LCRS) the largest length-scale at which filaments are statistically significant is between $70$ to $80 \\, h^{-1}$Mpc. Further we show that the average filamentarity of the galaxy distribution depends on various physical properties of galaxies such as luminosity, colour, morphology and star formation rate \\citep{pandey1, pandey3}. It would be interesting to know if the statistically significant length scale of filaments also depends on different galaxy properties. In the present work we study if the statistically significant length scale of filaments depends on luminosity, colour and morphology of galaxies.  Further it is also possible to measure the statistically significant length scale of filaments in mock galaxy distribution extracted from N-body simulation.  The \\lcdm model is currently believed to be the minimal model which is consistent with most cosmological data (\\citealt{efst}; \\citealt{perci}; \\citealt{teg1}; \\citealt{sperg}; \\citealt{sperg1}; \\citealt{komat}). It would be interesting to measure the statistically significant length scale of filaments in N-body simulations of $\\Lambda$CDM model and compare it with measured values in the SDSS \\citep{pandey}. The N-body simulations primarily predict the clustering of the dark matter. This should be contrasted with the fact surveys reveal only the bright side of the matter distribution. Galaxy formation is a complicated process and the exact relation between the distribution of the galaxies and the dark matter is not well understood. It is now generally accepted that the galaxies are a biased tracer of the dark matter distribution (e.g., \\citealt{kais}; \\citealt{mo}) and on large scales one expects the fluctuations in the galaxy and the dark matter distribution to be linearly related through the linear bias parameter $b$. Mock galaxy distributions with different bias can be simulated following this assumption. The average filamentarity of the simulated galaxy distribution is found to be very sensitive to the bias parameter \\citep{bharad3, pandey2}. It would be interesting to know how the statistically significant length scale of filaments depends on the bias parameter. In the present work we test if the statistically significant length scale of filaments depends on bias.  Earlier we find that the average filamentarity is sensitive to the area and galaxy number density of the samples \\citep{pandey1}(hereafter Paper II) and this statistics can be used for a meaningful comparison between two different galaxy samples only when they have the same volume (identical shape and size) and galaxy number density. In the present work we would also like to test if the statistically significant length scale of filaments depends on area and galaxy number density of the samples.  Finally we simulate segment Cox process \\citep{pons} with different segment lengths and test the efficiency of our method in measuring the statistically significant length scale of filaments in the mock samples drawn from these simulations.  A brief outline of our paper follows. Section 2 describes the data and method of analysis, our results and conclusions are presented in Section 3. ",
        "conclusions": "We use ``Shuffle'' to determine the statistically significant length scale of filaments in the seven SDSS strips in magnitude bin 1 and bin 2 (Table~1 and Table~2). We label bin 1 as Faint and bin 2 as Bright. The top two panels of Figure \\ref{fig:2} shows the results of ``Shuffle'' on strip 1 from bin 1 and bin 2 respectively. In top two panels of Figure \\ref{fig:2} we see that shuffling the data with $L=10$ causes a large drop in the average filamentarity ($F_2$). The average filamentarity increases as the value of $L$ increases and the value of average filamentarity of the unshuffled slice comes within $1 -\\sigma$ error bars when the data is shuffled with $L=90$ and $L=70$ for Bright and Faint galaxies respectively. Difference between the average filamentarity of shuffled and unshuffled slice is quantified by $\\chi^2/\\nu$ and the top two panels of Figure \\ref{fig:3} shows $\\chi^2/\\nu$ as a function of $L$. In top two panels of Figure \\ref{fig:3} we see that $\\chi^2/\\nu$ approaches 1 at $L=90$ for Bright and $L=70$ for Faint galaxies. So in strip 1 from bin1 and bin 2 the filaments are statistically significant upto length scales $80 \\, h^{-1} {\\rm Mpc}$ and $60 \\, h^{-1} {\\rm Mpc}$ respectively. The results from the other strips are not shown here. Averaging the results from all the nine strips we find that the filaments are statistically significant upto length scales $84 \\pm 16 \\, h^{-1} {\\rm Mpc}$ and $71 \\pm 16 \\, h^{-1} {\\rm Mpc}$ in bin 1 and bin 2 respectively.  The right and left middle panels of Figure \\ref{fig:2} shows the results of ``Shuffle'' for Red and Blue galaxies (Table~2) from strip 1 in bin 2. The right and left middle panels of Figure \\ref{fig:2} shows that a large drop in the average filamentarity is observed when the data is shuffled with $L=10$. In both the panels we see that shuffling the data with $L=80$ increase the average filamentarity and brings it back in agreement with the actual data. In two middle panels of Figure \\ref{fig:3} we see that in both class of galaxies $\\chi^2/\\nu$ approaches 1 at $L=80$ establishing that for both Red and Blue galaxies the filaments are statistically significant upto length scales $70 \\, h^{-1} {\\rm Mpc}$. Here we have not shown the results from the other strips. Combining results from all the nine strips we find that the filaments are statistically significant upto length scales $77 \\pm 10 \\, h^{-1} {\\rm Mpc}$ and $71 \\pm 10 \\, h^{-1} {\\rm Mpc}$ for Red and Blue galaxies respectively.  We show the results of ``Shuffle'' for the Early and Late type galaxies (Table~2) in the right and left bottom panels of Figure \\ref{fig:2}. Here also we see that the average filamentarity decreases when the data is shuffled with $L=10$ and it increases and matches with the actual data when the data is shuffled with $L=80$. The $\\chi^2/\\nu$ approaches 1 at $L=80$ for both Early and Late type galaxies as shown in two bottom panels of Figure \\ref{fig:3}. This establishes that for both Early and Late type galaxies the filaments are statistically significant upto length scales $70 \\, h^{-1} {\\rm Mpc}$ and not beyond. We have shown here only the results for strip 1 in bin 2. Combining the results from all the strips the filaments are found to be statistically significant upto length scales $75 \\pm 10  \\, h^{-1} {\\rm Mpc}$ and $74 \\pm 17  \\, h^{-1} {\\rm Mpc}$ for Early and Late type galaxies respectively.  It is to be noted in middle and bottom two panels of Figure \\ref{fig:2} that the Red and Early type galaxies have higher degree of filamentarity as compared to Blue and Late type galaxies at smaller values of filling factor. We reported this effect in \\citet{pandey1}. It is interesting to note that although the average filamentarity ($F_2$) depends on luminosity, colour and morphology of galaxies, the statistically significant length scale of filaments does not depend on these galaxy properties.  In Figure \\ref{fig:4} we show the results of ``Shuffle'' in simulated SDSS NGP strips with different biases. The top left panel show the result for the dark matter distribution from $\\Lambda$CDM N-body simulations. In this case galaxies are assumed to exactly trace the dark matter and the bias parameter $b=1$. We see that shuffling the simulated data with $L=10$ causes a large drop in the average filamentarity of the simulated galaxy distribution. With increasing L the value of average filamentarity of the shuffled slice slowly approaches the values corresponding to the unshuflled data. It is found that at $L=90$ the two are within $1 -\\sigma$ error bars. The top right panel shows the result for b=1.2. We see a very similar result and the unshuflled and shuffled data agrees well when the data is shuffled with $L=90$. In top left and right panels of Figure \\ref{fig:5} we see that $\\chi^2/\\nu$ approaches 1 at $L=90$ for both $b=1$ and $b=1.2$ and hence filaments are statistically significant upto length scales $80 \\, h^{-1} {\\rm Mpc}$ for both cases. The results are shown for one single mock SDSS strip. Averaging the results from all the nine strips it is found that for b=1 and b=1.2 filaments are statistically significant upto length scales $93  \\pm 16  \\, h^{-1} {\\rm Mpc}$ and $92  \\pm 15  \\, h^{-1} {\\rm Mpc}$ respectively.  The bottom left panel of Figure \\ref{fig:4} shows the result of shuffling on simulated SDSS strips for a bias b=0.8. As usual shuffling the data with $L=10$ causes a drop in average filamentarity of the shuffled slice but the drop is smaller as compared to other bias values. As L increases the average filamentarity slowly grows and approaches the values corresponding to the unshuffled data at $L=90$. In the bottom right panel of Figure \\ref{fig:4} results are shown for a high bias b=1.5. Shuffling shows similar results for b=1.5 but a relatively large drop is seen when the data is shuffled with $L=10$. With increasing L the average filamentarity grows relatively faster as compared to b=0.8 and finally the average filamentarity in the shuffled data levels up with the unshuffled data when $L=80$ is used for shuffling. In the bottom left and right panels of Figure \\ref{fig:5} it is found that $\\chi^2/\\nu$ approaches 1 at $L=90$ and $L=80$ indicating that filaments are statistically significant upto length scales $80 \\, h^{-1} {\\rm Mpc}$ and $70 \\, h^{-1} {\\rm Mpc}$ for b=0.8 and b=1.5 respectively. These are the results shown for a single mock SDSS strip and by combining the results from the other strips we get that the filaments are statistically significant upto length scales $91 \\pm 13  \\, h^{-1} {\\rm Mpc}$ and $90 \\pm 15 \\, h^{-1} {\\rm Mpc}$ for $b=0.8$ and $b=1.5$ respectively.  \\begin{figure} \\rotatebox{-90}{\\scalebox{.4}{\\includegraphics{plot5.ps}}} \\caption{This shows $\\chi^2/\\nu$ at different shuffling lengths for a simulated SDSS NGP strip with different bias values as indicated in each panel. The black dotted line indicates $\\chi^2/\\nu=1$.} \\label{fig:5} \\end{figure}  \\begin{figure} \\rotatebox{-90}{\\scalebox{.4}{\\includegraphics{plot6.ps}}} \\caption{This shows the statistically significant length scale of filaments as a function of bias b. The error bars are $1-\\sigma$ error bars measured from nine mock galaxy catalogues.} \\label{fig:6} \\end{figure}  \\begin{figure} \\rotatebox{-90}{\\scalebox{.4}{\\includegraphics{plot7.ps}}} \\caption{This shows the statistically significant length scale of filaments as a function area of the samples described in Table~3. The error bars are $1 - \\sigma$ error bars measured from nine mock galaxy catalogues. } \\label{fig:7} \\end{figure}  \\begin{figure} \\rotatebox{-90}{\\scalebox{.4}{\\includegraphics{plot8.ps}}} \\caption{This shows the statistically significant length scale of filaments as a function galaxy number density of the samples described in Table~3. The error bars are $1 - \\sigma$ error bars measured from nine mock galaxy catalogues. } \\label{fig:8} \\end{figure}  It is to be noted in Figure \\ref{fig:4} that as the bias $b$ is increased the average filamentarity ($F_2$) increases at smaller $FF$. We reported this effect in \\citet{bharad3}. But interestingly the shuffling length at which the avergae filamentarity of the unshuflled and shuffled data agrees does not depend on bias. Figure \\ref{fig:6} shows the statistically significant length scale of filaments as a function of bias. This Figure shows that the statistically significant length scale of filaments is nearly independent of bias. The mean value is $~90 \\, h^{-1} {\\rm Mpc} $ for the bias values $b=0.8-1.5$. It tends to decrease for a high value of bias $b=1.8$. This could be possibly related to the fact that at very high values of bias we preferentially identify only very high density regions which are mostly related to cluster like structures rather than filaments.  Earlier in Paper I we established that the filaments are statistically significant upto length scales of $80 \\, h^{-1} {\\rm Mpc}$. This measured value lies well within $1-\\sigma$ error bars of the measured values for mock galaxy distribution from $\\Lambda$CDM N-body simulations. We conclude that the $\\Lambda$CDM is consistent with SDSS observations.  The galaxy samples constructed over different magnitude ranges have different area and galaxy number density (Table~2). Earlier analysis in Paper II shows that the average filamentarity of the galaxy distribution depends on the area and galaxy number density of the galaxy samples and filamentarity between different galaxy samples can only be compared when the samples have same area, geometry and number density of galaxies. So it is important to know if the statistically significant length scale of filaments depends on these factors before we can compare it's measured values among various galaxy samples having different areas and galaxy number densities.  We prepare two different set of mock galaxy samples namely area 1 and area 2 (Table~3) which have the same galaxy number density as the SDSS NGP strip (area 3) but cover different redshift ranges and hence have different areas as listed in Table~3. We extract three mock galaxy samples from each of the $\\Lambda$CDM dark matter simulation for each set of galaxy samples area 1 and area 2. We determine the statistically significant length scale of filaments for these galaxy samples. The analysis for the mock SDSS NGP strip which we name as area 3 in Table~3 are already done. We plot the statistically significant length scale of filaments as a function of the area of the galaxy samples in Figure \\ref{fig:7}. From this figure we see that the samples of smaller area tend to have a smaller length scale of filaments. But this trend is weak and the size of the error bars are quite large. We conclude that the statistically significant length scale of filaments are nearly independent of the area of the samples.  We next prepare two different set of mock galaxy samples namely density 1 and density 3 (Table~3) which have the same area as the SDSS NGP strip (density 2) but have different number densities listed in Table~3. We extract three mock galaxy samples from each of the $\\Lambda$CDM dark matter simulation for each set of galaxy samples density 1 and density 3. We determine the statistically significant length scale of filaments for these galaxy samples. The analysis for the mock SDSS NGP strip which we name as density 2 in Table~3 are already done. We plot the statistically significant length scale of filaments as a function of the number density of galaxies of the samples in Figure \\ref{fig:8}. We see a trend of low density samples to have a smaller length scale of filaments. But the error bars are quite large and the dependence is weak. We conclude that the statistically significant length scale of filaments weakly depends on the number density of galaxies. In a recent work with SDSS LRG samples \\citep{pandey4} having much larger area and lower density we find that the statistically significant length scale of filaments does not change much as compared to it's value measured in SDSS Main galaxy samples. Thus the statistically significant length scale of filaments is a more robust statistics than the average filamentarity. It does not depend on properties (size and number density) of galaxy samples , and physical properties of galaxies. It emerges as a robust measure of the large scale structures. The statistically significant length scale of filaments could be also thought of as an indicator of the scale beyond which the galaxy distribution become homogeneous because it tells us the length scale beyond which no coherent structures exist in the galaxy distribution. Multifractal analysis in SDSS DR1 \\citep{yadav} and recently in SDSS DR6 \\citep{sarkar} show that the galaxy distribution becomes homogeneous at a length scale between $60$ to $70 h^{-1} \\, {\\rm Mpc}$. \\citet{yadav} considered the transition to homogeneity in $\\Lambda$CDM simulations with different bias $b=1,1.6,2$ and find that irrespective of the bias values the distribution become homogeneous at a length scale between $60$ to $70 h^{-1} \\, {\\rm Mpc}$. These results are consistent with our findings.  Finally we test the reliability of our method by applying it to different controlled samples (Table~4) of 3D segment Cox process. We use segments of two different length $l=20\\, h^{-1} {\\rm Mpc}$ and $l=80\\, h^{-1} {\\rm Mpc}$ for simulating the 3D segment Cox process. Two different sets of 2D mock samples drawn from these 3D simulations (Table~4) are separately analyzed. The results are shown in top two panels of Figure \\ref{fig:9} and Figure \\ref{fig:10}. We see in top left panel of Figure \\ref{fig:9} that filamentarity decrease by a very small amount when the mock samples having $l=20\\, h^{-1} {\\rm Mpc}$ are shuffled with $L=10$ and there are virtually no drop in filamentarity when the data is shuffled with $L=20$. The top right panel of the same figure shows the result of ``Shuffle'' on cotrolled samples of segment Cox process with segment length $l=80\\, h^{-1} {\\rm Mpc}$. We see a significant drop in the average filamentarity when the data is shuffled with $L=10$ and the average filamentarity finally saturates to the values corresponding to the unshuffled data when $L=60$ is used for shuffling. The $\\chi^2/\\nu$ as a function of $L$ for these two cases are shown in top two panels of Figure \\ref{fig:10}. We see that $\\chi^2/\\nu$ approaches 1 at $L=20$ and $L=60$ in these two cases respectively. This indicates that for these 2D mock samples having $l=20\\, h^{-1} {\\rm Mpc}$ and $l=80\\, h^{-1} {\\rm Mpc}$, the filaments are found to be statistically significant upto length scales $10 \\, h^{-1} {\\rm Mpc}$ and $50 \\, h^{-1} {\\rm Mpc}$ respectively. Here the results are shown only for a single simulated mock sample. By averaging the results from all the nine simulated mock samples we find that in the above two cases the filaments are statistically significant upto length scales $18 \\pm 7 \\, h^{-1} {\\rm Mpc}$ and $52 \\pm 11 \\, h^{-1} {\\rm Mpc}$ respectively. We note that the values of statistically significant length scale of filaments recovered by our method for straight segments simulated here are a little shorter but close to the original values of $l$ used in the simulations. The simulations are carried out in 3D boxes and our mock samples are the 2D projections of nearly two dimensional samples drawn from these boxes. This possibly destroys and shortens a lot of segments and the effects are expected to be more prominent for longer segments as reflected in our results. This effect indicates that also for the 2D real galaxy samples from SDSS which are constructed out of a 3D galaxy distribution, Shuffle possibly underestimates the values of statistically significant length scale of the real 3D filaments.  The effects of slicing and projections can be avoided if we simulate the segment Cox process in 2D as in this case the segments lie on the same plane (Figure \\ref{fig:11}). We apply Shuffle to a set of mock 2D strips constructed out of simulations of 2D segment Cox process (Table~5). We use two different segment lengths $l=20\\, h^{-1} {\\rm Mpc}$ and $l=80\\, h^{-1} {\\rm Mpc}$ to simulate the 2D segment Cox process. We use the same segment lengths as used in the simulations of 3D segment Cox process so as to test the effect of slicing and projection. We show the results in bottom two panels of Figure \\ref{fig:9} and Figure \\ref{fig:10}. Bottom left panel of Figure \\ref{fig:9} shows that when the mock samples with $l=20\\, h^{-1} {\\rm Mpc}$ are shuffled with $L=10$ the average filamentarity of the shuffled data falls below that of the unshuffled data. The average filamentarity of the shuffled data saturates to the unshuffled data when the data is shuffled with $L=40$. The bottom right panel of Figure \\ref{fig:9} shows a significant drop in the average filamentarity when the mock samples with $l=80\\, h^{-1} {\\rm Mpc}$ are shuffled with $L=10$. The average filamentarity of the shuffled data increases slowly with increasing shuffling length and finally saturates to the unshuffled data at shuffling length $L=90$. We see in bottom two panels of Figure \\ref{fig:10} the $\\chi^2/\\nu$ as a function of $L$ approaches 1 at $L=40$ and $L=90$ in these two cases respectively. This indicates that for these controlled samples of 2D segment Cox process having $l=20\\, h^{-1} {\\rm Mpc}$ and $l=80\\, h^{-1} {\\rm Mpc}$, the filaments are statistically significant upto length scales $30 \\, h^{-1} {\\rm Mpc}$ and $80 \\, h^{-1} {\\rm Mpc}$ respectively. We have shown the results for a single simulated strip in these plots. We combine the results from all the nine simulated strips and find that the filaments are statistically significant upto length scales $25 \\pm 8 \\, h^{-1} {\\rm Mpc}$ and $82 \\pm 24 \\, h^{-1} {\\rm Mpc}$ for simulations of 2D segment Cox process with segment length $l=20\\, h^{-1} {\\rm Mpc}$ and $l=80\\, h^{-1} {\\rm Mpc}$ respectively.  We note that our method can successfully recover better the length of inputted segments when the mock samples are drawn from simulations of 2D segment Cox process where the effects of slicing and projections are absent. There would be another effect due to the boundaries of the samples which again shortens the length of some segments near the boundaries but this effect is not so important in our analysis as the extent of the samples are much larger than the length of the inputted segments. This analysis suggets that the actual length scale of the real 3D filaments can be recovered with Shuffle if the analysis is extended in 3D. We propose to take up this in a future work.  In conclusion we note that although the average filamentarity of the galaxy distribution depends on different galaxy properties, the statistically significant length scale of filaments does not depend on luminosity, colour and morphology of galaxies. We find that the measured values of the statistically significant length scale of filaments in SDSS are consistent with that measured from $\\Lambda$CDM simulations. We considered the $\\Lambda$CDM model with different values of bias and find that the statistically significant length scale of filaments is nearly independent of bias. Different class of galaxies are differently biased with respect to underlying dark matter distribution and this result is possibly related to the fact that statistically significant length scale of filaments is nearly the same for different class of galaxies. Unlike the average filamentarity, the statistically significant length scale of filaments is nearly independent of the area and galaxy number density of the samples. This establishes the statistically significant length scale of filaments as a robust statistics of galaxy distribution.  \\begin{figure} \\rotatebox{0}{\\scalebox{.6}{\\includegraphics{plot11.ps}}} \\caption{ This shows a simulation of 2D segment Cox process with two different segment length $l=20\\, h^{-1} {\\rm Mpc}$ and $l=80\\, h^{-1} {\\rm Mpc}$ (as indicated in each panel) over a region which has identical geometry to the SDSS NGP equatorial strip. These simulated strips have same values of $l$ and $\\lambda_{l}$ listed in Table~5 but different values of $\\lambda_{s}$ are chosen so as to lower the density of the segments and make them visibly distinguishable in this plot.} \\label{fig:11} \\end{figure}  \\begin{figure} \\rotatebox{-90}{\\scalebox{.4}{\\includegraphics{plot9.ps}}} \\caption{ This shows the Average Filamentarity as a function of Filling Factor ($FF$) for a mock SDSS NGP strip constructed out of 3D simulation of segment Cox process (Table~4) and 2D simulation of segment Cox process (Table~5) together with the results for the shuffled data for two values of $L$ shown in the figure. The 2D simulations of segment Cox process contain all the segments in the same plane and there are no effects because of the slicing process or the projections which arises when samples are constructed out of 3D simulations of segment Cox process. We see in this plot that ``Shuffle\" is able to recover the length of the longer segments more efficiently when samples are constructed out of 2D simulations instead of 3D simulations of segment Cox process. } \\label{fig:9} \\end{figure}  \\begin{figure} \\rotatebox{-90}{\\scalebox{.4}{\\includegraphics{plot10.ps}}} \\caption{ This shows $\\chi^2/\\nu$ at different shuffling lengths ($L$) for a mock SDSS NGP strip constructed out of 3D segment Cox process and 2D segment Cox process simulated with two different segment length $l=20\\, h^{-1} {\\rm Mpc}$ and $l=80\\, h^{-1} {\\rm Mpc}$. The black dotted line indicates $\\chi^2/\\nu=1$.} \\label{fig:10} \\end{figure}"
    },
    "0910/0910.4579.txt": {
        "abstract": "{The dynamical ejection of stars from star clusters affects the shape of the stellar mass function (MF) in these clusters, because the escape probability of a star depends on its mass. This is found in $N$-body simulations and has been approximated in analytical cluster models by fitting the evolution of the MF. Both approaches are naturally restricted to the set of boundary conditions for which the simulations were performed.} {The objective of this paper is to provide and to apply a simple physical model for the evolution of the MF in star clusters for a large range of the parameter space. It should also offer a new perspective on the results from $N$-body simulations.} {A simple, physically self-contained model for the evolution of the stellar MF in star clusters is derived from the basic principles of two-body encounters and energy considerations. It is independent of the adopted mass loss rate or initial mass function (IMF), and contains stellar evolution, stellar remnant retention, dynamical dissolution in a tidal field, and mass segregation.} {The MF evolution in star clusters depends on the disruption time, remnant retention fraction, initial-final stellar mass relation, and IMF. Low-mass stars are preferentially ejected after $t\\sim 400$~Myr. Before that time, masses around 15---20\\% of the maximum stellar mass are lost due to their rapid two-body relaxation with the massive stars that still exist at young ages. The degree of low-mass star depletion grows for increasing disruption times, but can be quenched when a large fraction of massive remnants is retained. The highly depleted MFs of certain Galactic globular clusters are explained by the enhanced low-mass star depletion that occurs for low remnant retention fractions. Unless the retention fraction is exceptionally large, dynamical evolution always decreases the mass-to-light ratio. The retention of black holes reduces the fraction of the cluster mass in remnants because white dwarfs and neutron stars have masses that are efficiently ejected by black holes.} {The modeled evolution of the MF is consistent with $N$-body simulations when adopting identical boundary conditions. However, it is found that the results from $N$-body simulations only hold for their specific boundary conditions and should not be generalised to all clusters. It is concluded that the model provides an efficient method do understand the evolution of the stellar MF in star clusters under widely varying conditions.} % ",
        "introduction": "\\label{sec:intro} The evaporation of star clusters is known to change the shape of the underlying stellar mass function\\footnote{Hereafter, `mass function' is referred to as `MF'.} \\citep{henon69,chernoff90,vesperini97b,takahashi00,portegieszwart01,baumgardt03}. This phenomenon has been used to explain the observed MFs in globular clusters \\citep{richer91,demarchi07,demarchi07b}, which are flatter than typical initial mass functions \\citep[IMFs, e.g.][]{salpeter55,kroupa01}. In addition, the effect of a changing MF on cluster photometry has been investigated \\citep{lamers06,kruijssen08,anders09}. This has been shown to explain the low mass-to-light ratios of globular clusters \\citep{kruijssen08b,kruijssen09} and to have a pronounced effect on the inferred globular cluster mass function \\citep{kruijssen09b}.  The existing parameterised cluster models that incorporate a description of low-mass star depletion are restricted by the physically self-contained models on which they are based. Some studies \\citep{lamers06,kruijssen08} assume an increasing lower stellar mass limit to account for the evolving MF, others \\citep{anders09} fit a changing MF slope to $N$-body simulations. In both cases, the models are accurate for a certain range of boundary conditions, but they do not include a physical model and are therefore lacking flexibility. While $N$-body simulations do include the appropriate physics, they are very time-consuming. As a result, only a limited number of clusters can be simulated and the applicability of the simulations is thus restricted to the specific set of boundary conditions for which they have been run.  It would be desirable to obtain a simple physical model for the evolution of the MF, which would have a short runtime and could be used independently of $N$-body simulations. Forty years ago, a pioneering first approach to such a model was made by \\citet{henon69}, who considered the stellar mass-dependent escape rate of stars from star clusters. However, the applicability of his model was limited due to a number of assumptions that influenced the results. First of all, \\citet{henon69} assumed that the clusters exist in isolation and neglected the tidal field. As a consequence, the ejection of a star could only occur by a single, close encounter and the repeated effect of two-body relaxation was not included. Secondly, the {distribution of stars was independent of stellar mass, i.e.} mass segregation was not included. Both mass segregation and the influence of a tidal field are observed in real clusters, and can be expected to affect the evolution of the MF.  The aim of this paper is to derive a physical description of the evolution of the stellar MF in star clusters, alleviating the assumptions that were made by \\citet{henon69}. This should explain the results found in $N$-body simulations and observations, while providing the required flexibility to explore the properties of star clusters with simple, physically self-contained models. The outline of this paper is as follows. In Sect.~\\ref{sec:model}, total mass evolution of star clusters is discussed. A recipe for the evolution of the MF is derived in Sect.~\\ref{sec:mf}, covering stellar evolution, the retain of stellar remnants, dynamical dissolution and mass segregation. The model is compared to $N$-body simulations in Sect.~\\ref{sec:comp}. In Sect.~\\ref{sec:results}, the model is applied to assess the evolution of the MF for different disruption times and remnant retention fractions. The consequences for other cluster properties are also considered. This paper is concluded with a discussion of the results and their implications. ",
        "conclusions": "\\label{sec:disc} The results of this paper show that the stellar MFs in star clusters differ strongly from their initial forms due to dynamical cluster evolution. The specific kinds of these differences depend on the properties of the star clusters and their tidal environment, most importantly on the disruption time, remnant retention fraction, and IMF.\\footnote{Although not specifically shown in this paper (but not surprisingly), the differences also depend on the initial-final stellar mass relation.}  A physical model for the evolution of the stellar MF is presented in which two-body relaxation leads to a stellar mass dependence of the ejection rate. For any particular stellar mass, the ejection rate is determined by the typical proximity of that mass to the escape energy and by the timescale on which the two-body relaxation with the other stars takes place. Combined with a prescription for stellar evolution, stellar remnant production, and remnant retention using kick velocity dispersions, this provides a description for the total evolution of the MF. {\\it This description is independent of the adopted total mass evolution}. {The model shows that the slope of the mass function is a possible indicator for the mass fraction that has been lost due to dissolution, provided that the IMF does not vary and the remnant retention fraction has been fairly similar for young globular clusters.\\footnote{Any variability of the retention fraction would induce substantial scatter, see Sect.~\\ref{sec:remn} and Fig.~\\ref{fig:alphatdis}.}}  For the exact same initial conditions, the model shows excellent agreement with $N$-body simulations of the evolving MF by \\citet{baumgardt03}. However, an important advantage of the presented model compared to the (more accurate) $N$-body simulations is its short runtime and corresponding flexibility. It can be easily applied to compute the evolution of clusters for a large range of initial conditions. The results can then be used to identify interesting cases for more detailed and less simplified calculations with $N$-body or Monte Carlo models.  {The most important simplification of the model is neglecting the effect of binary encounters on the stellar mass dependence of the ejection rate. To incorporate binaries, a conclusive census of the binary population in star clusters would be required, which is not yet available. Nonetheless, it is possible to make a qualitative estimate for the effect of binaries. The encounter rate of binaries would typically be higher than that of individual stars, because the cross section of binaries is larger. This would increase the relative escape rate at the stellar mass for which the binary fraction\\footnote{The fraction of stars residing in binary or multiple systems.} peaks. This binary fraction is found to increase with primary mass \\citep[see e.g.][]{kouwenhoven09}. Because massive stars are removed by stellar evolution, this implies that the binary fraction decreases with age, which is in agreement with the low binary fraction observed in globular clusters \\citep[$\\sim 2\\%$, e.g.][]{richer04}. The effect of binaries on the evolution of the mass function would thus be most notable if the majority of the dynamical mass loss occurs at ages $< 50$~Myr (the typical lifetime of an $8~\\msun$ star), in which case it would somewhat enhance the relative escape rate of the most massive stars.}  The model is applied to investigate the influence of the disruption time and remnant retention on the evolution of the MF and integrated photometric properties of star clusters. For total disruption times $\\tdis<1$~Gyr, the modeled relative ejection rate is highest at a certain `sweet spot' mass that is typically 15---20\\% of the mass of the most massive objects in the cluster. For longer lifetimes, the evolution of the MF is dominated by low-mass star depletion, unless the retention fraction of massive stellar remnants is larger than 0.25. Only in the particular case of such a high retention fraction, the $M/L$ ratio is increased by dynamical evolution when the cluster approaches total disruption. In all other scenarios, the $M/L$ ratio decreases because the most massive (luminous) stars are kept.\\footnote{This process differs from a possible variability of the proportionality between the velocity dispersion and the cluster mass, which concerns a much shorter timescale \\citep[e.g.][]{boily09}.} When defining the slope of the MF in the range 30---80\\% of the maximum stellar mass, this gives a clear relation between the MF slope and the $M/L$ ratio. For slopes that are defined in fixed mass ranges, there is not necessarily a correlation between slope and $M/L$ ratio if $\\tdis<1$~Gyr. In clusters with a longer total disruption time, both quantities are related. Dynamical cluster evolution is found to induce some reddening of the integrated cluster colours, amounting up to 0.1---0.2~mag in $V-I$ for total disruption times $\\tdis<1.5$~Gyr. The fraction of the cluster mass that is constituted by remnants surprisingly decreases if more black holes are retained, because the black holes preferentially eject bodies around the masses of white dwarfs and neutron stars, which contain most of the total remnant mass.  Contrary to what is suggested by other studies \\citep[e.g.][]{baumgardt03,anders09}, the evolution of the MF is not homologous. The reason that these studies concluded that its evolution is very similar for all clusters (also see Figs.~\\ref{fig:comp} and~\\ref{fig:comp2}), is that they assumed that all remnants were retained. It is illustrated in Fig.~\\ref{fig:mf2} that the differences between clusters with dissimilar disruption times disappear when the retention fraction increases. For realistic retention fractions, differences do arise. If two clusters with different initial masses have the same total disruption time, their MF evolution will be dissimilar due to their different remnant retention fractions and the impact of the retained remnants on the dynamical cluster evolution. Alternatively, if two clusters have equal initial masses but different total disruption times, for instance due to differences in their galactic location or environment, their MF evolution will be dissimilar due to the dynamical impact of the evolution of the maximum stellar mass.  The larger variation of MF evolution that is found with presented model may also be able to explain observations of globular clusters in which the MF cannot be characterised by a single power law \\citep{demarchi00}. If the evolution of the MF were homologous, these features would likely be primordial \\citep{baumgardt03}, but this is not necessarily the case when using realistic retention fractions. Most other differences between the results presented in Sect.~\\ref{sec:results} and those from \\citet{baumgardt03} are also due to their assumption of full remnant retention. For example, their $M/L$ ratio evolution shows a smaller decrease than in Fig.~\\ref{fig:ml1}. This is explained in Fig.~\\ref{fig:ml2}, where it is shown that dynamical evolution reduces the $M/L$ ratio by a smaller amount if the retention fraction is larger.  {Studies on the fractal nature of cluster formation show that star clusters are initially substructured \\citep{elmegreen00,bonnell03}. Even though this substructure is typically erased on a crossing time, it can induce primordial mass segregation in star clusters \\citep{mcmillan07,allison09}.} The influence of primordial mass segregation on the evolution of the MF has recently been investigated by \\citet{baumgardt08b} and \\citet{vesperini09}. While \\citet{baumgardt08b} do not include stellar evolution and concentrate on two-body relaxation, \\citet{vesperini09} do include stellar evolution. They show that for some degrees of primordial mass segregation, the mass loss by stellar evolution can induce additional dynamical mass loss that strongly decreases the total disruption time. For clusters that survive for a Hubble time, the MF evolution in the case of primordial mass segregation is very similar to an initially unsegregated cluster. \\citet{vesperini09} conclude that the evolution of the MF is only affected by primordial mass segregation for clusters in which the total disruption time is sufficiently decreased by the induced mass loss. In that case, the slope of the MF remains much closer to its initial value than it would in clusters without primordial mass segregation. Their conclusion is consistent with the model presented in this paper, because the evolution of the MF is determined by the most massive stars at the time when the largest mass loss occurs (see Figs.~\\ref{fig:chinorm} and~\\ref{fig:mf1}). This induced mass loss enters the model in terms of the absolute mass loss rate in Eq.~\\ref{eq:dmdt}, not in the stellar mass-dependent escape rate per unit mass loss rate of Eq.~\\ref{eq:chi}.  A change in total mass loss rate is not the only consequence of primordial mass segregation. \\citet{baumgardt08b} have shown that low-mass star depletion is enhanced for clusters without stellar evolution that are primordially mass-segregated. This occurs because energy equipartition is reached on a shorter timescale and because of their use of a fixed ($m_{\\rm max}=1.2~\\msun$) maximum stellar mass. As a result, there are no massive bodies to increase the ejection rate of intermediate mass stars (see Fig.~\\ref{fig:maxchinorm}), implying that only the low-mass stars are preferentially lost. In the present paper, mass segregation is assumed to arise dynamically, but the model could in principle be adapted to cover primordial mass segregation by setting $c_2=0$ and adjusting $c_1$ to the initial velocity distribution until it is erased by dynamical evolution (see Eq.~\\ref{eq:tseg}), after which the values from Sect.~\\ref{sec:mf} can be used.\\footnote{As explained in Sect~\\ref{sec:mf}, $c_1$ represents the ratio of the mean speed squared to the central escape velocity squared that depends on the degree of mass segregation (and thus on the IMF). On the other hand, $c_2$ is a proportionality constant in the expression for the onset of the stellar mass-dependent ejection of stars, which depends on the concentration or King parameter.} This does not necessarily yield enhanced low-mass star depletion for clusters with a complete IMF (including masses $m>1.2~\\msun$) because of the presence of massive stars or remnants.  \\begin{figure}[t] \\resizebox{\\hsize}{!}{\\includegraphics{alphatdis.ps}} \\caption[]{\\label{fig:alphatdis} MF slope versus remaining lifetime (assuming a globular cluster age of 12~Gyr). Diamonds represent the observed values from \\citet{demarchi07}, with typical errors as shown by the error bar in the lower right corner. The remaining lifetimes are taken from \\citet{baumgardt08b}. Dotted curves represent the model evolutionary tracks of clusters with $\\log{(M_{\\rm i}/\\msun)}=\\{6,6.25,6.5,6.75,7\\}$ from Sect.~\\ref{sec:remn} with $\\{\\sigma_{\\rm kick,wd},\\sigma_{\\rm kick,ns},\\sigma_{\\rm kick,bh}\\}=\\{4,100,200\\}$~km~s$^{-1}$, corresponding to $\\{f_{\\rm ret,wd},f_{\\rm ret,ns},f_{\\rm ret,bh}\\}=\\{0.983,0.022,0.003\\}$ for a $10^6~\\msun$ cluster. The solid line connects the present-day locations of the modeled clusters in the diagram (crosses), while the dashed line represents the same relation for $\\sigma_{\\rm kick,bh}=40$~km~s$^{-1}$ ($f_{\\rm ret,bh}=0.219$ for a $10^6~\\msun$ cluster). The dash-dotted line shows the homologous cluster evolution from \\citet{baumgardt03}. } \\end{figure} The presented model can be applied to the MFs of Galactic globular clusters that are observed by \\citet{demarchi07}. These MFs are more strongly depleted than is found in the $N$-body simulations by \\citet{baumgardt03}, which has been attributed to primordial mass segregation \\citep{baumgardt08b}. However, the observations can also very accurately be explained with the realistic remnant retention fractions that are used in the present paper. This is shown in Fig.~\\ref{fig:alphatdis}, where the observed MF slopes and remaining lifetimes of the globular clusters from \\citet{demarchi07} are compared with the globular cluster-like models from Sect.~\\ref{sec:remn} {($t_0=1$~Myr)}. The models are in much better agreement with the data than the $N$-body runs with complete remnant retention from \\citet{baumgardt03}. Deviations to other values of $\\alpha$ can occur due to variations in disruption time and remnant retention fractions, as is also shown in Fig.~\\ref{fig:alphatdis}. For example, a variation of the remnant kick velocity with metallicity in combination with the known variation of the disruption time \\citep[see, e.g.][]{kruijssen09,kruijssen09b} should be sufficient to cover the observed scatter.  {The above line of reasoning provides an explanation for the the depleted MFs in Fig.~\\ref{fig:alphatdis} that is consistent with the simulations by \\citet{vesperini09}, who showed that the effects of primordial mass segregation are in fact suppressed in long-lived clusters due to the expansion caused by stellar evolution. This increases the relaxation time and yields an evolution of the MF that is very similar to the initially unsegregated scenario, indicating that primordial mass segregation is not a likely explanation for strongly depleted MFs.} Observations of the remnant composition of these globular clusters could reveal {a definitive answer as to} whether the depleted MFs are explained by primordial mass segregation or by dynamical evolution with a realistic remnant retention fraction.  Dynamical cluster evolution does not appear to have a large effect on the colours of old (globular) clusters. The only way in which the colours could be affected beyond typical observational errors, is if globular clusters have lost substantial fractions of their masses during the first $\\sim$~Gyr after their formation. In that case, the dynamical evolution of the stellar MF in globular clusters may have implications for studies of colour bimodality \\citep[e.g.][]{larsen01} or the blue tilt \\citep[e.g.][]{harris06}. It could then also possibly explain the trend of increasing $V-K$ colour with decreasing $M/L_V$ ratio found by \\citet{strader09} for globular clusters in M31, because quickly dissolving clusters generally become redder and have reduced $M/L$ ratios. More research is needed to determine the role of the changing MF in the above properties of globular cluster systems.  It can be concluded that the evolution of the stellar MF in star clusters is not as similar for all clusters as previously thought. Its precise evolution is determined by cluster characteristics like the disruption time, the remnant retention fraction, initial-final stellar mass relation, and the IMF. In order to decipher the evolution of observed star clusters, it is essential to record these characteristics and to relate them to possible scenarios for the internal evolution of clusters. That way, observables like the slope of the MF, the $M/L$ ratio, the broadband colours, and the mass fraction in remnants can be better understood."
    },
    "0910/0910.5134_arXiv.txt": {
        "abstract": "We propose a new scenario for the early universe where there is a smooth transition between an early de Sitter-like phase and a radiation dominated era. In this model, the matter content is modelled  by a new type of generalised Chaplygin gas \\cite{chaplygin} for the early universe, with an underlying scalar field description. We study the gravitational waves generated by the quantum fluctuations. In particular, we calculate the gravitational wave power spectrum, as it would be measured today, following  the method of the Bogoliubov coefficients. We show that the high frequencies region of the spectrum depends strongly  on one of the parameters of the model. On the other hand, we use the number of e-folds, along with the power spectra and spectral index of the scalar perturbations, to constrain the  model observationally. ",
        "introduction": "\\label{sec1}  The inflationary paradigm is the most consistent scenario to explain the origin of the large scale structure (LSS) of the universe, as well as the anisotropies in the cosmic microwave background radiation (CMBR) \\cite{LiddleLyth,Kinney:2009vz}. Originally, an early inflationary era in the universe  was invoked as a mechanism to solve several shortcomings of the big bang theory \\cite{inflation}. Afterwards, it was realised  that an early inflationary era of the universe generates density perturbations as seeds for the structure we see nowadays. In addition, a background of stochastic primordial gravitational waves (GW) is also predicted and it is  originated from the vacuum fluctuations. The latter (GW) if ever detected in the future would provide an amusing source of information about the early universe. On the other hand, and most importantly, the predictions of the inflationary theory have been corroborated by several cosmological observations like the recent WMAP5 data \\cite{WMAP}.   In this letter we propose a phenomenological model for the inflationary era and the subsequent radiation dominated epoch of the universe. In particular, we suggest a way to extend the framework of the generalised Chaplygin gas (GCG) to the first stages of the universe\\footnote{An alternative inflationary model inspired on the Chaplygin gas was proposed in Ref.~\\cite{Bertolami:2006zg} where only the inflationary era is accounted for.}. The GCG has attracted a lot of attention over the last years  \\cite{chaplygin,chaplygin2,gorini,Barreiro:2008pn,GCGsingularity}. It was initially introduced  in cosmology as a mean to unify the dark sectors of the universe, i.e. dark energy and dark matter, and has been contrasted with several cosmological observations \\cite{Barreiro:2008pn}. On the other hand, it has been as well shown that some GCG models can be an excellent framework to analyse dark energy related singularities \\cite{chaplygin2,GCGsingularity}.   In this work, we investigate the possible imprints in the power spectrum of the stochastic background of GW, for the model we propose for the transition from the inflationary era to the radiation dominated epoch\\footnote{In Ref.~\\cite{{Fabris:2004dp}}, the spectrum of GW for a universe whose dark energy component corresponds to a Chaplygin gas was analysed. This scenario is different from ours.}. This transition is far from being well known and it can be of some interest to explore the signatures associated with it, shedding light on the inflationary model behind the accelerated expansion in the early universe.   We calculate the GW production using the Bogoliubov coefficients which obey a set of differential equations \\cite{parker,Starobinsky,allen,mendes-henriques-moorhouse}. This method has advantages over the frequently used sudden transition approximation, because, associated with it, there is always an overproduction of gravitons of large frequencies, which is avoided in a natural way by the use of the continuous Bogoliubov coefficients \\cite{mendes-henriques-moorhouse}. This is also a very practical method to calculate the full spectrum, from the very low frequencies corresponding to the present cosmological horizon $10^{-17}$Hz, till those large, kHz up to GHz, frequencies associated with the transition between the inflationary  and the radiation-dominated universe.   The present letter is organised as follows. In section \\ref{sec2}  we present our model based on a \\textit{new type of generalised Chaplygin gas for the early universe}. This fluid can be as well represented by a minimally coupled scalar field as shown in section \\ref{sec3}. In section \\ref{sec4}, we constrain our model observationally. In section \\ref{sec5}, we summarise the  methodology used to obtain the spectrum of the stochastic GW which is based on the Bogoliubov coefficients. We present in section \\ref{sec:NumericalSimulations} the spectrum of the GW for the model introduced in sections \\ref{sec2} and \\ref{sec3}. Finally, in section \\ref{sec7} we present our conclusions. ",
        "conclusions": "\\label{sec7} \\label{sec:Conclusions}  In this work, we investigate the production of GWs in a new modified generalised Chaplygin gas for the early universe (cf. Secs \\ref{sec2} and \\ref{sec3}). Through recent measurements of CMBR/LSS, the parameters of the model have been constrained.   We have used the method of the continuous Bogoliubov coefficients to calculate the GW energy spectrum for different scales of energy and values of the $\\alpha$-parameter of the model.  Besides the fact that the model suffers a small deviation from the preferred values of the spectral index $n_s$, the obtained spectra reveal a consistent picture corresponding to a smooth transition between the inflationary and the radiation dominated epochs of the universe, and constitute a significant imprint of this modified generalised Chaplygin gas.  Finally, the strong variation of the high frequency range, is directly related with the decrease in the maximum of the potential $a^{\\prime\\prime}/a$ (see Eq.~(\\ref{6})) for $|\\alpha|$ approaching $1$. This feature implies strong limits to the maximum frequency allowed in our model and which will be within the reach of future gravitational-waves detectors like BBO and DECIGO \\cite{Lidsey97}, for the KHz range of frequencies. In fact, for the most consistent values of the $\\alpha$-parameter, the spectrum shows a frequency as low as in the Hz region.  Concluding, these results show that, for this model, the Hz-KHz frequency range comes directly from the transition between the inflationary regime and the radiation era, and addresses important issues about the limits of the GW energy spectrum.  Last but not least, we would like to highlight that one of the merits of the model we have presented for the transition from the inflationary era to the radiation dominated epoch is its relative simplicity."
    },
    "0910/0910.2652_arXiv.txt": {
        "abstract": "We present a Chandra X-ray observation of the very high energy $\\gamma$-ray source HESS$~$J1640-465. We identify a point source surrounded by a diffuse emission that fills the extended object previously detected by XMM Newton at the centroid of the HESS source, within the shell of the radio supernova remnant (SNR) G338.3-0.0. The morphology of the diffuse emission strongly resembles that of a pulsar wind nebula (PWN) and extends asymmetrically to the South-West of a point-source presented as a potential pulsar. The spectrum of the putative pulsar and compact nebula are well-characterized by an absorbed power-law model which, for a reasonable $N_{\\rm H}$ value of $14\\times 10^{22} \\rm cm^{-2}$, exhibit an index of 1.1 and 2.5 respectively, typical of Vela-like PWNe. We demonstrate that, given the H$~$I absorption features observed along the line of sight, the SNR and the H$~$II surrounding region are probably connected and lie between 8 kpc and 13 kpc. The resulting age of the system is between 10 and 30 kyr. For a 10 kpc distance (also consistent with the X-ray absorption) the 2-10 keV X-ray luminosities of the putative pulsar and nebula are $L_{\\rm PSR} \\sim 1.3 \\times 10^{33}  d_{10 \\rm ~kpc}^{2} \\rm erg.s^{-1}$ and $L_{\\rm PWN} \\sim 3.9 \\times 10^{33}  d_{10}^{2} \\rm erg.s^{-1}$ ($d_{10} = d / 10{\\rm~kpc}$). Both the flux ratio of $L_{\\rm PWN}/L_{\\rm PSR} \\sim 3.4$  and the total luminosity of this system predict a pulsar spin-down power around $\\dot{E} \\sim 4 \\times 10^{36} \\rm erg~s^{-1}$. We finally consider several reasons for the asymmetries observed in the PWN morphology and discuss the potential association with the HESS source in term of a time-dependent one-zone leptonic model. ",
        "introduction": "During 2004-2006, H.E.S.S. (High energy stereoscopic system) performed a survey of the inner part of the Galaxy where its excellent capability provided a breakthrough in the field of pulsar wind nebulae (PWN) study: for the first time the morphological structure of many middle-age PWNe were resolved in the $\\gamma$-ray band. A detailed spectral and morphological analysis of the archetypal example of this population, PSR$~$B1823-13 and its associated TeV nebula HESS$~$J1825-137 revealed for the first time in $\\gamma$-rays a steepening of the energy spectrum with increasing distance from the pulsar, probably due to the synchrotron radiative cooling of the Inverse-Compton (IC) $\\gamma$-ray emitting electrons during their propagation (Aharonian et al. 2006b). In this scenario, the different sizes observed in X-rays and very high energy (VHE) $\\gamma$-rays can be explained by the different cooling timescales for the radiating electron populations (de Jager et al. 2005). However, the question of the unusual large size of this PWN candidate compared to what the dynamical simulations predict still remains (de Jager et al. 2005; de Jager $\\&$ Djannati 2008) and is difficult to solve. The confusion is even aggravated because of the non-detection of the associated supernova shell, like it is the case for most of the other middle-aged TeV PWN candidates (Gallant et al. 2007).  HESS$~$J1640-465 is one of the sources discovered by HESS during this Galactic survey in 2004-2006 (Aharonian et al. 2006a). This source is marginally extended with profile close to a Gaussian shape. The spectrum is well fitted by a simple power-law with an index of $\\sim$ 2.4 and a total integrated flux above 200 GeV of $2.2 \\times 10^{-11} {\\rm ~ erg ~ cm}^{-2}{\\rm ~ s}^{-1}$. As it is shown in Figure 1, the source is spatially coincident with the SNR G338.3-0.0, known from radio observations as a $8'$ diameter broken shell with a 7 Jy flux at 1GHz (Whiteoak $\\&$ Green 1996). A radio bridge extending to the north of the system is visible in the MOST data and coincident with the bright H$~$II region G338.4+0.0 (Whiteoak $\\&$ Green 1996), possibly connected to the shell. The fact that the TeV source is center filled and doesn't match the shell has favoured the hypothesis that the emission is produced by a PWN so far. This field has been successively observed by ASCA (Sugizaki et al. 2001) and the Swift X-ray telescope (Landi et al. 2006); both detected a highly absorbed source inside the shell, with compatible positions and non thermal spectra.\\\\ In 2006, a dedicated XMM-Newton observation of the field of view revealed an extended X-ray object at the position ($16^{\\rm h}40^{\\rm m}45^{\\rm s}.4$, $-46^{\\rm o}31'31''$(J2000))(Funk et al. 2007). This extended source (XMMUJ164045.4-463131) exhibits a strongly absorbed spectrum with a potential non-thermal spectrum (index $\\sim 1.7$) centered about the VHE source position. However, because of the low number of counts of these observations the spectrum was poorly constrained, and the insufficient angular resolution was not able to resolve any compact nebula, nor a pulsar which could confirm the hypothesis of a PWN emission.  We report here on the first Chandra observation of this field. We present a detailed data analysis and discuss the physical implications of this first case of a complete middle-age composite system with VHE PWN emission. ",
        "conclusions": "This high-resolution Chandra ACIS observation of HESS$~$J1640-165 inside the shell G338.3-0.0, has allowed us to spatially resolve a PWN and its putative pulsar, and perform a spectral analysis. We found that the spectra are both highly absorbed and non-thermal, and the indices and luminosities resemble those of typical Vela-like pulsar systems. Using H$~$I absorption features, we were able to constrain the distance of the shell between 8.5 and 13 kpc, this distance range being compatible with the large $N_{\\rm H}$ value derived from X-ray spectral analysis, we concluded that it is likely associated with the shell G338.3-0.0. We finally discussed a scenario in which the $\\gamma$-rays originate from the IC scattering and X-ray from the synchrotron emission of the electron populations in the nebula, and show that the VHE source HESS$~$J1640-165 is likely a PWN, probably associated with a Vela-like pulsar of $\\dot{E} =  4 \\times 10^{36} \\rm erg.s^{-1}$. If this interpretation is true, this object is a nice example of a middle-aged composite system where the total PWN extension (traced by the TeV emission) fills almost 75 percent of the associated SNR shell area, which is quite surprising given that dynamical simulations currently predict a 25 percent. Future search for similar cases will be useful in order to confirm or not the generality of this phenomenon and study its causes in detail. In this context, the detection of the associated pulsars will be very helpful. A next step of this study could be the high resolution radio imaging of the source and the search for pulsed emission originating from this putative pulsar position."
    },
    "0910/0910.0723.txt": {
        "abstract": "{The present number of Galactic Open Clusters that have high resolution abundance determinations, not only of [Fe/H], but also of other key elements, is largely insufficient to enable a clear modeling of the Galactic Disk chemical evolution. } {To increase the number of Galactic Open Clusters with high quality measurements. } {We obtained high resolution (R$\\sim$30\\,000), high quality (S/N$\\sim$50-100 per pixel), echelle spectra with the fiber spectrograph FOCES, at Calar Alto, Spain, for three red clump stars in each of five Open Clusters. We used the classical Equivalent Width analysis method to obtain accurate abundances of sixteen elements: Al, Ba, Ca, Co, Cr, Fe, La, Mg, Na, Nd, Ni, Sc, Si, Ti, V, Y. We also derived the oxygen abundance through spectral synthesis of the 6300~\\AA\\  forbidden line. } {Three of the clusters were never studied previously with high resolution spectroscopy: we found [Fe/H]=+0.03$\\pm$0.02 ($\\pm$0.10)~dex for Cr~110; [Fe/H]=+0.01$\\pm$0.05 ($\\pm$0.10)~dex for NGC~2099 (M~37) and [Fe/H]=--0.05$\\pm$0.03 ($\\pm$0.10)~dex for NGC~2420. This last finding is higher than typical recent literature estimates by 0.2--0.3~dex approximately and in better agreement with Galactic trends. For the remaining clusters, we find: [Fe/H]=+0.05$\\pm$0.02 ($\\pm$0.10)~dex for M~67 and [Fe/H]=+0.04$\\pm$0.07 ($\\pm$0.10)~dex for NGC~7789 . Accurate (to $\\sim$0.5~km~s$^{-1}$) radial velocities were measured for all targets, and we provide the first high resolution based velocity estimate for Cr~110, $<$$V_r$$>$=41.0$\\pm$3.8~km~s$^{-1}$. } {With our analysis of the new clusters Cr~110, NGC~2099 and NGC~2420, we increase the sample of clusters with high resolution based abundances by 5\\%. All our programme stars show abundance patterns which are typical of open clusters, very close to solar with few exceptions. This is true for all the iron-peak and s-process elements considered, and no significant $\\alpha$-enhancement is found. Also, no significant sign of (anti-)correlations for Na, Al, Mg and O abundances is found. If anticorrelations are present, the involved spreads must be $<$0.2~dex. We then compile high resolution data of 57 OC from the literature and we find a gradient of [Fe/H] with Galactocentric Radius of --0.06$\\pm$0.02~dex~kpc$^{-1}$, in agreement with past work and with Cepheids and B stars in the same range. A change of slope is seen outside $R_{\\rm{GC}}$=12~kpc and [$\\alpha$/Fe] shows a tendency of increasing with $R_{\\rm{GC}}$. We also confirm the absence of a significant Age-Metallicity relation, finding slopes of --2.6$\\pm$1.1~10$^{-11}$~dex~Gyr$^{-1}$ and 1.1$\\pm$5.0~10$^{-11}$~dex~Gyr$^{-1}$ for [Fe/H] and [$\\alpha$/Fe] respectively. } ",
        "introduction": "Open clusters (hereafter OC) are the ideal {\\em test particles} in the study of the Galactic disk, providing chemical and kinematical information in different locations and at different times. Compared to field stars, they have the obvious advantage of being coeval groups of stars, at the same distance and with a homogeneous composition. Therefore, their properties can be determined with smaller uncertainties. Several attempts have been done in the past to derive two fundamental relations using OC: the {\\em metallicity gradient} along the disk and the {\\em age-metallicity relation} (hereafter AMR) of the disk \\citep[e.g.][]{jan79,pan80,tw97,friel02,che03,sal04}, but they were hampered by the lack of large and homogeneous high quality datasets.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% Observing Logs \\begin{table*} \\begin{minipage}[t]{17.5cm} \\caption{Observing Logs and Programme Stars Information.} \\label{obslog} \\centering \\renewcommand{\\footnoterule}{} \\begin{tabular}{l l c c c c c c c c c c c } \\hline\\hline Cluster      & Star & $\\alpha_{J2000}$ & $\\delta_{J2000}$ &  B   &  V   &I$_{C}$&  R   & K$_S$& $n_{exp}$ & $t_{exp}^{(tot)}$ & S/N($\\simeq$6000\\AA) \\\\ &      & (hrs)            & (deg)            &(mag) &(mag) & (mag) &(mag) &(mag) &           & (sec)             &               \\\\ \\hline Cr~110\\footnote{Star names and I$_C$ \\& R magnitudes from \\citet{daw98}; coordinates and B \\& V magnitudes from \\citet{bra03}; K$_S$ magnitudes from 2MASS} & 2108 & 06:38:52.5       & +02:01:58.4\t  &14.79 &13.35 &  ---  & ---  & 9.76 & 6\t  & 16200             & 70     \\\\ & 2129 & 06:38:41.1       & +02:01:05.5\t  &15.00 &13.66 & 12.17 &12.94 &10.29 & 7 \t  & 18900             & 70     \\\\ & 3144 & 06:38:30.3       & +02:03:03.0\t  &14.80 &13.49 & 12.04 &12.72 &10.19 & 6\t  & 16195             & 65     \\\\ NGC~2099 (M~37)\\footnote{Star names from \\citet{vanz21}; coordinates from \\citet{kiss01}; B \\& V magnitudes from \\citet{kalirai01}; I$_C$ magnitudes from \\citet{nila02}; K$_S$ magnitudes from 2MASS} & 067  & 05:52:16.6       & +32:34:45.6\t  &12.38 &11.12 &  9.87 & ---  & 8.17 & 3\t  &  3600             & 95     \\\\ & 148  & 05:52:08.1       & +32:30:33.1\t  &12.36 &11.09 &  ---  & ---  & 8.05 & 3\t  &  3600             & 105    \\\\ & 508  & 05:52:33.2       & +32:27:43.5\t  &12.24 &10.98 &  ---  & ---  & 7.92 & 3\t  &  3900             & 85     \\\\ NGC~2420\\footnote{Star names from \\citet{can70}; coordinates from \\citet{stet00} and \\citet{las90}; B \\& V magnitudes from \\citet{tw90}; I$_C$ \\& R magnitudes from \\citet{stet00}; K$_S$ magnitudes from 2MASS} & 041  & 07:38:06.2       & +21:36:54.7\t  &13.75 &12.67 & 11.61 &12.13 &10.13 & 5\t  &  9000             & 70     \\\\ & 076  & 07:38:15.5       & +21:38:01.8\t  &13.65 &12.66 & 11.65 &12.14 &10.31 & 5 \t  &  9000             & 75     \\\\ & 174  & 07:38:26.9       & +21:38:24.8\t  &13.41 &12.40 &  ---  & ---  & 9.98 & 5\t  &  9000             & 60     \\\\ NGC~2682 (M~67)\\footnote{Star names from \\citet{fag06}; coordinates from \\citet{fan96}; B, V \\& I$_C$ magnitudes from \\citet{san04}; B magnitude for star 286 and R magnitudes from \\citet{jan84}; K$_S$ magnitudes from 2MASS} & 0141 & 08:51:22.8       & +11:48:01.7\t  &11.59 &10.48 &  9.40 & 9.92 & 7.92 & 3 \t  &  2700             & 85     \\\\ & 0223 & 08:51:43.9       & +11:56:42.3\t  &11.68 &10.58 &  9.50 &10.02 & 8.00 & 3 \t  &  2700             & 85     \\\\ & 0286 & 08:52:18.6       & +11:44:26.3\t  &11.53 &10.47 &  9.43 & 9.93 & 7.92 & 3 \t  &  2700             & 105    \\\\ NGC~7789\\footnote{Star names and V \\& I$_c$ magnitudes from \\citet{gim98b}; J1950 coordinates from \\citet{kun23}, precessed to J2000; B magnitudes from \\citet{moc99}; K$_S$ magnitudes from 2MASS} & 5237 & 23:56:50.6       & +56:49:20.9\t  &13.92 &12.81 & 11.52 & ---  & 9.89 & 5 \t  &  9000            &  70   \\\\ & 7840 & 23:57:19.3       & +56:40:51.5\t  &14.03 &12.82 & 11.49 & ---  & 9.83 & 6\t  &  9000            &  75   \\\\ & 8556 & 23:57:27.6       & +56:45:39.2\t  &14.18 &12.97 & 11.65 & ---  &10.03 & 3\t  &  5400            &  45   \\\\ \\hline \\end{tabular} \\end{minipage} \\end{table*} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% Observing Logs  In particular, the lack of a {\\em metallicity scale} extending to solar metallicity with comparable precision to the lower metallicity regimes \\citep[i.e.,][]{zin84,car97} represents the main problem from the point of view of {\\em (i)} the study of the Galactic disk; {\\em (ii)} tests of stellar evolution models for younger and more metal-rich {\\em simple stellar populations} and {\\em (iii)} the use of those stellar populations as templates for extragalactic studies of population synthesis. Of the $\\sim$1700 known OC \\citep[][and updates]{dia02}, only a subset of $\\sim$140, i.e., 8\\% of the total, possesses some metallicity determination. Most of these have been obtained through different photometric studies in the Washington \\citep[e.g.][]{gei91,gei92}, DDO \\citep[e.g.][]{cla99}, Str\\\"omgren \\citep[e.g.][]{bru99,twa03}, UBV \\citep[e.g.][]{cam85} and IR \\citep[e.g.][]{tie97} photometric systems and passbands, giving often rise to considerable differences with those obtained from spectroscopy \\citep[see][and re\\-fe\\-ren\\-ces therein]{gra00}. In a much smaller number of clusters, abundances have been derived from low-resolution spectroscopy \\citep[e.g.,][]{ricardo,war08}, with admirable attempts to obtain large and homogeneous datasets \\citep[see][]{friel93,friel02}, in spite of the non-negligible uncertainties involved in the procedure.  A few research groups (see Section~\\ref{sec-disc} for more details) are presently obtaining high quality spectra and are producing more precise abundance determinations. The study of elements other than the iron-peak ones (such as $\\alpha$, s- and r-process, light elements), allows one to put more constraints on the sites of production of those elements (SNe Ia, SNe II, giants and supergiants, Wolf-Rayet stars), and therefore on their production timescales. These are fundamental ingredients for the chemical evolution modeling of the Galactic Disk \\citep{tos82,chi01,col09}.  For these reasons, we obtained high resolution spectra for a sample of poorly studied old OC. We present here the detailed abundance analysis of five clusters observed during our first run at Calar Alto. Observations and data reductions are described in Section~\\ref{sec-obs}; the linelist and equivalent width measurements are detailed in Section~\\ref{sec-li} while the abundance analysis methods and results are presented in Section~\\ref{sec-abo}; abundance results are then discussed and compared with literature results in Sections~\\ref{sec-lit}, \\ref{sec-disc} and \\ref{sec-trend}; finally, we summarize our results and draw our conclusions in Section~\\ref{sec-sum}.   %__________________________________________________________________ ",
        "conclusions": "\\label{sec-sum}  We have analyzed high resolution spectra of three red clump giants in five OC, three of them lacking any previous high resolution based chemical analysis. Given the paucity of literature data, such a small sample is enough to increase the whole body of high resolution data for OC by $\\simeq$5\\%. To compare our results with the literature, we have compiled chemical abundances based on high resolution data of 57 clusters from the literature. Given the recent and fast progress in the field, this sample is $\\sim$50\\% larger than previous literature compilations \\citep[e.g.,][]{friel02,sestito,mag09}. The main results drawn by the analysis of our five clusters are:  \\begin{itemize} \\item{We provide the first high resolution based abundance analysis of Cr~110 ([Fe/H]=+0.03$\\pm$0.02 ($\\pm$0.10)~dex), NGC~2099 ([Fe/H]=+0.01$\\pm$0.05 ($\\pm$0.10)~dex) and NGC~2420 ([Fe/H]=--0.05$\\pm$0.03 ($\\pm$0.10)~dex), which only had low resolution determinations and the R$\\simeq$16000 analysis by \\citet{smi87}; our new determination of the metallicity of NGC~2420 puts this cluster in much better agreement with the global Galactic trends;} \\item{The abundances found for NGC~7789 ([Fe/H]= +0.04$\\pm$0.07 ($\\pm$0.10)~dex) and M~67 ([Fe/H]= +0.05$\\pm$0.02 ($\\pm$0.10)~dex) are in good agreement with past high resolution studies;} \\item{We provide the first high resolution based radial velocity determination for Cr~110 ($<V_r>$=41.0$\\pm$3.8~km~s$^{-1}$);} \\item{We found that all our abundance ratios, with few exceptions generally explained with technical details of the analysis procedure, are near-solar, as is typical for OC with similar properties; we found solar ratios also for Na, that is generally found overabundant, and for O, which is generally found underabundant;} \\item{We do not find any significant sign of anti-correlation (or correlation) among Na, Al, Mg and O, in general agreement with past and recent results, and we can say that if such correlations indeed are present in OC, they must be much less extended than in Globular Clusters, amounting to no more than 0.2--0.3~dex at most;} \\end{itemize}  With our compilation of literature data, we also could examine global Galactic trends, that are extremely useful to construct chemical evolution models for the Galaxy in general and the Galactic Thin Disc in particular. For the metallicity gradient we found a slope of  --0.06$\\pm$0.02~dex~kpc$^{-1}$ considering only the high resolution data within $R_{\\rm{GC}}$=12~kpc. Our compilation contains data, including our own determinations, that fill the small gap around 9$<$R$_{\\rm{GC}}<$12 and point towards a gentle and continuous decrease, rather than a two step drop such as in \\citet{tw97} or a steep slope such as in \\citet{sestito}. We do find a flattening at $R_{\\rm{GC}}>$12~kpc and a hint of an increasing [$\\alpha$/Fe] towards the outer Disk. Concerning the AMR, we do not find any strong evidence for it, and we just note some very mild trends. If an AMR is indeed present among OC, it must be very weak and we provide upper limits to its slope both in [Fe/H] and [$\\alpha$/Fe].  %______________________________________________________________"
    },
    "0910/0910.5843_arXiv.txt": {
        "abstract": "We test the theoretical prediction that the straightest dust lanes in bars are found in strongly barred galaxies, or more specifically, that the degree of curvature of the dust lanes is inversely proportional to the strength of the bar. The test used archival images of barred galaxies for which a reliable non-axisymmetric torque parameter ($Q_{\\rm b}$) and the radius at which $Q_{\\rm b}$ has been measured ($r(Q_{\\rm b})$) have been published in the literature. Our results confirm the theoretical prediction but show a large spread that cannot be accounted for by measurement errors. We simulate 238 galaxies with different bar and bulge parameters in order to investigate the origin of the spread in the dust lane curvature versus $Q_{\\rm b}$ relation. From these simulations, we conclude that the spread is greatly reduced when describing the bar strength as a linear combination of the bar parameters $Q_{\\rm b}$ and the quotient of the major and minor axis of the bar, $a/b$. Thus we conclude that the dust lane curvature is predominantly determined by the parameters of the bar. ",
        "introduction": "Dust lanes in galactic bars have been observed in nearby galaxies for a long time. Photographic surveys from the beginning of 20$^{\\rm th}$ century already allowed observation of these features. Pease (1917) describes the prominent and nearly straight dust lanes in NGC~5383 as a `dark streak'. The same paper presents images of NGC~1068 in which a dust lane is clearly visible. The dusty nature of these features was identified at a  much later stage. Sandage (1961) writes in his atlas that 'one of the major characteristics of the SBb(s) [galaxies] is the presence of two dust lanes leaving the nucleus, one on each side of the bar and extending into the spiral arms'.  Dust lanes have been recognised as related to shocks in the gas flow in barred galaxies since Prendergast (1962). Such shocks are usually found in the leading edge of the bar, roughly parallel to its major axis. The location of these shocks also corresponds to the location of areas of high shear, which prevents star formation (Athanassoula 1992; see Zurita et al.~2004 for a graphic illustration in the case of NGC~1530). Simulations show that the dust lanes are nearly straight and near the center of the bar when the galaxy has no inner Lindblad resonance (ILR), which is known to be quite a rare occurrence. ILRs cause the dust lanes to be offset from the bar major axis, and to be curved with the concavity pointing to the bar major axis (Athanassoula 1992). Simulations by Patsis \\& Athanassoula (2000) show that the higher the gas sound speed, the smaller is the offset between the dust lane and the major axis of the bar. No dust lanes are expected in nuclear bars (Shlosman \\& Heller 2002).  Athanassoula (1992) predicted from simulations that dust lanes would have a greater curvature in weaker bars and that dust lanes would be nearly straight in strong bars. This effect has been empirically verified using a set of images of nine barred spiral galaxies by Knapen et al.~(2002). The aim of the present Letter is to improve the statistics compared to the latter study, and to test definitively the prediction from Athanassoula (1992). Because dust lanes can be observed so easily they offer a fundamental handle on the underlying dynamics of the galaxy and the bar, yet they have hardly been studied observationally. A basic further aim of the present work is thus to explore in a statistical fashion the dependence of dust lane shape on the dynamical properties of the galaxy which hosts them. ",
        "conclusions": "Using 55 galaxies with published measurements of the bar strength indicator $Q_{\\rm b}$, and for which we could obtain accurate measurements of the dust lane curvature, $\\Delta\\alpha$, we confirm the theoretical prediction made by Athanassoula (1992) that the dust lane curvature anticorrelates with the bar strength. Strong bars thus host straight dust lanes, while weaker bars can host more curved dust lanes. We do find that the anticorrelation has a large spread, but from a set of 238 numerical simulations of barred galaxies we show that this spread can be greatly reduced by using an appropriate linear combination of bar parameters ($Q_{\\rm b}$ and the axis ratio $a/b$)."
    },
    "0910/0910.1698_arXiv.txt": {
        "abstract": "{}{We use the IBIS/ISGRI telescope on-board \\emph{INTEGRAL} to measure the position of the centroid of the 20-200 keV emission from the Crab region.}{We find that the astrometry of the IBIS telescope is affected by the temperature of the IBIS mask during the observation. After correcting for this effect, we show that the systematic errors on the astrometry of the telescope are of the order of 0.5 arcsec. In the case of the Crab nebula and several other bright sources, the very large number of photons renders the level of statistical uncertainty in the centroid smaller or comparable to this value.}{We find that the centroid of the Crab nebula in hard X-rays (20-40 keV) is shifted by 8.0 arcsec with respect to the Crab pulsar in the direction of the X-ray centroid of the nebula. A similar shift is also found at higher energies (40-100 and 100-200 keV). We observe a trend of decreasing shift with energy, which can be explained by an increase in the pulsed fraction. To differentiate between the contribution of the pulsar and the nebula, we divide our data into an on-pulse and off-pulse sample. Surprisingly, the nebular emission (i.e., off-pulse) is located significantly away from the X-ray centroid of the nebula.}{In all 3 energy bands (20-40, 40-100 and 100-200 keV), we find that the centroid of the nebula is significantly offset from the predicted position. We interpret this shift in terms of a cut-off in the electron spectrum in the outer regions of the nebula, which is probably the origin of the observed spectral break around 100 keV. From a simple spherically-symmetric model for the nebula, we estimate that the electrons in the external regions of the torus ($d\\sim0.35$ pc from the pulsar) reach a maximal energy slightly below $10^{14}$ eV.} ",
        "introduction": "The Crab nebula is the brightest astrophysical source in the $\\gamma$-ray sky ($E>30$ keV). It is the prototypical pulsar-wind nebula (PWN, see e.g., \\citet{gaensler} for a review), where material ejected from the central pulsar at different epochs interacts, producing strong shocks that accelerate electrons up to energies $\\sim10^{15}$ eV \\citep{kennel}. The synchrotron emission produced by the large population of non-thermal electrons in a magnetic field $B\\sim0.3$ mG \\citep{marsden} is detected over more than 10 orders of magnitude in the electromagnetic spectrum from radio to $\\gamma$-rays \\citep{atoyan,volpi,hesterreview}.  In soft X-rays, the size of the nebula is $\\sim$2 arcmin \\citep{weisskopf}. It consists of a characteristic jet+torus structure, probably corresponding to relativistic outflows along the rotation axis and the equator of the pulsar. High-resolution \\emph{Chandra} observations also revealed the presence of an inner ring, probably associated with the conversion of the relativistic pulsar wind into a synchrotron-emitting plasma. The measured photon index decreases with radius from $\\Gamma\\sim1.6$ around the pulsar down to $\\Gamma\\sim3.0$ in the outer parts of the X-ray image \\citep{mori,willingale}, which implies that the most energetic electrons are constantly injected by the pulsar in the inner regions of the nebula. The steeper photon index in the outer regions can be explained by synchrotron cooling. This interpretation is supported by the observation of rapid X-ray variability in the inner ring \\citep{hester}, which is probably due to the presence of strong quasi-stationary shocks. The pulsar itself exhibits a hard spectrum $\\Gamma\\sim1.6$ and a period $P\\sim30$ msec with a double-peak profile.  In the hard X-ray/soft $\\gamma$-ray range, early HEAO A-4 observations inferred a photon index of $\\Gamma=2.15$ \\citep{jung}. Around 100 keV a spectral break was detected, and above this energy a photon index of $\\sim$2.5 was found. This measurement agrees with the steeper spectral indices measured at higher energies by COMPTEL \\citep{strong} and EGRET \\citep{nolan}. Therefore, observations of the Crab complex around the break energy by modern experiments are important to constrain particle acceleration models. Recently, a polarized $\\gamma$-ray signal from the Crab was measured by the SPI \\citep{dean} and IBIS \\citep{forot} instruments on \\intgr. The polarization was found to be co-aligned with the spin orientation, thus indicating a possible association of the polarized signal with the inner jet. The broad-band coverage of \\intgr\\ also permitted a study of the evolution of the pulse profiles as a function of energy \\citep{mineo}. A phase-revolved analysis detected a significantly harder signal ($\\Gamma\\sim1.8$) during the interpulse phase compared to the peaks ($\\Gamma\\sim2.2$). However, the results presented there only considered the emission from the pulsar, and no information was given about the spectrum of the nebula.  In this paper, we use the coded-aperture IBIS telescope \\citep{Ubeetal03} on-board \\intgr\\ \\citep{Winetal03} to measure the position of the hard X-ray/soft $\\gamma$-ray centroid of the Crab pulsar/nebula complex. In Sect. \\ref{astrom}, we analyse the astrometry of the IBIS telescope and show that the point-source localization accuracy of the instrument depends on the temperature of the mask. We correct for this effect and assess the level of systematic uncertainties in the astrometry of IBIS. In Sect. \\ref{results}, we report precise measurements of the position of the Crab pulsar/nebula and present our results. The implications of these findings are discussed in Sect. \\ref{disc}. ",
        "conclusions": "We have used the IBIS/ISGRI instrument on-board \\intgr\\ to measure the centroid of the Crab pulsar/nebula complex in hard X-rays/soft $\\gamma$-rays with unprecedented accuracy, with the aim of investigating the behaviour of the relativistic electron population around the break energy ($E_{\\mbox{\\tiny{break}}}\\sim100$ keV). Based on our understanding of the dependence of the astrometry of IBIS on the temperature of the mask, we have shown in Sect. \\ref{syst} that despite its moderate angular resolution (12 arcmin FWHM), in the case of sources detected with high $S/N$ IBIS/ISGRI can measure the position of astrophysical sources with an accuracy $\\sim0.5$ arcsec. Applying this method to the Crab, we found that the emission is offset significantly with respect to the position of the pulsar, by $\\sim8\"$ in the 20-40 keV band. We were also able to measure the position of the on- and off-pulse phases independently. All the measured positions are found along the line connecting the pulsar to the centroid of the X-ray emission in the torus, which corresponds roughly to the jet line (see Sect. \\ref{interp}).  Performing phase-resolved imaging (see Sect. \\ref{relapos}), we found that the on-pulse emission (where the Crab pulsar accounts for an important fraction of the flux) depends on energy: the on-pulse shift with respect to the pulsar position decreases by $\\sim30\\%$ between 20 and 200 keV. This effect can be easily explained by an increase in the pulsed fraction (see Sect. \\ref{pulsfrac}).  More interestingly, we find that during the off-pulse phase (where the emission from the pulsar is negligible) the centroid of the source is significantly offset from the position predicted from \\cra\\ data (see Sect. \\ref{disc}). This indicates that the morphology of the source changes around the break energy. Consistently, we show that a cut-off in the electron spectrum in the outer regions of the X-ray nebula, caused by synchrotron cooling and in agreement with the predictions of \\citet{atoyan}, can reproduce the observed shift.  The centroid of the nebular emission around the break energy coincides with the inner ring, which is interpreted as the location to be the wind termination shock \\citep{weisskopf}. Therefore, our data imply that above the break energy only the strong shock regions are responsible for the emission. Comparing our results with the predictions of a simple model for the spectral evolution of the nebula, we find that the electrons in the outer regions of the torus ($d\\gtrsim0.35$ pc away from the pulsar) probably reach a maximal energy close to $10^{14}$ eV. This result agrees with theoretical studies carried out using MHD simulations \\citep{volpi}. In the near future, because of a higher angular resolution by a factor $\\sim20$ \\emph{NuSTAR} will be able to resolve the Crab pulsar/nebula complex up to $\\sim80$ keV, which will allow us to probe the electron spectrum in the different regions of the nebula close to the break energy."
    },
    "0910/0910.4242_arXiv.txt": {
        "abstract": "CCD {\\it UBVRI} photometry is presented for type IIb SN 2008ax for about 320 days. The photometric behavior is typical for core-collapse SNe with low amount of hydrogen. The main photometric parameters are derived and the comparison with SNe of similar types is reported. Preliminary modeling is carried out, and the results are compared to the observed light curves. The main parameters of the hydrodynamical models are close to those used for SN IIb 1993J. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.4418_arXiv.txt": {
        "abstract": "Approximately 30\\% of luminous red giants exhibit a Long Secondary Period (LSP) of variation in their light curves, in addition to a shorter primary period of oscillation.  The cause of the LSP has so far defied explanation: leading possibilities are binarity and a nonradial mode of oscillation.  Here, large samples of red giants in the Large Magellanic Cloud both with and without LSPs are examined for evidence of an 8 or 24 $\\mu$m mid-IR excess caused by circumstellar dust.  It is found that stars with LSPs show a significant mid-IR excess compared to stars without LSPs.  Furthermore, the near-IR $J$-$K$ color seems unaffected by the presence of the 24 $\\mu$m excess.  These findings indicate that LSPs cause mass ejection from red giants and that the lost mass and circumstellar dust is most likely in either a clumpy or a disk-like configuration.  The underlying cause of the LSP and the mass ejection remains unknown. ",
        "introduction": "When red giants become more luminous than $L \\sim 1000$\\,\\lsun, they begin to vary with periods of variation which fall on six physically distinct period-luminosity sequences\\footnote{On these sequences, stars on the first ascent giant branch (FGB) have slightly longer periods than those on the asymptotic giant branch (AGB) at the same luminosity \\citep{ita04,sosetal07}.  Some authors designate the corresponding FGB and AGB sequences as distinct sequences.} \\citep{woo99,ita04,sosetal07,fra08}.  As well as these six sequences for the most luminous red giants, there is another sequence (known as sequence-E) extending to lower luminosities and which is known to consist of close binary systems exhibiting ellipsoidal light variations \\citep{woo99,sos04}.  Among the six luminous sequences all but the longest in period can be explained by radial pulsation in the fundamental and overtone modes \\citep{woo99}.  The longest period sequence, sequence-D, consists of stars that exhibit variations with a short period, usually corresponding to radial pulsation in a low overtone mode, as well as a Long Secondary Period (LSP).  It is the LSPs that make up sequence-D in the period-luminosity diagram for variable red giants.  Approximately 30\\% of all luminous red giants have LSPs \\citep{woo99,per03,sosetal07,fra08} so the LSP phenomenon is very common in the late evolution of low mass stars.  \\citet{sosetal07} suggest that LSPs can be seen at very low amplitude in up to 50\\% of luminous red giants.  The LSPs also seem to occur in red supergiants \\citep{sto71}.  Since the period of the LSP is approximately 4 times longer than the period of the normal-mode radial fundamental mode, the LSP can not be due to normal-mode radial pulsation.  The most favoured explanations for the origin of LSPs are binarity and nonradial g-modes \\citep{woo99,hin02,wok04,der06,sos07,hin09}, but these and other explanations for the LSPs all have significant problems \\citep{wok04,nic09}.  For example, stars exhibiting LSPs have velocity curves which generally are of similar shape and, in a binary model, this would indicate a favoured value for the angle of periastron of the orbit rather than the expected uniform distribution \\citep{nic09}.  In addition, the full amplitudes of the velocity curves are closely concentrated around 3.5 km s$^{-1}$ \\citep{nic09} which implies the unlikely situation wherein all these stars have companions of a similar mass of $\\sim$0.09 M$_{\\odot}$.  The problem with the nonradial g-mode explanation for the LSPs is that g-modes, which have the correct periods, only have substantial amplitudes in radiative regions and red giants have convective envelopes with only a thin radiative layer on top \\citep{wok04}.  Currently, no satisfactory explanation for the LSPs exists.  Several of the suggested explanations for the LSPs in red giants involve circumstellar dust.  The change in visible light associated with the LSP is typically of a rather irregular nature, and the light variation can be large, up to a factor of two.  In the binary model, it has been suggested that the light change could be due to a large dust cloud orbiting with the companion and obscuring the red giant once per orbit \\citep{woo99}.  Another possible explanation for the LSPs \\citep{woo99} involves semi-periodic dust ejection events such as those predicted by theoretical models of AGB stars \\citep{win94,hof95}.  Confirmation that dust was involved with the LSP phenomenon would be obtained if these stars showed a mid-infrared flux excess due to absorption of stellar light by circumstellar dust followed by re-radiation in the mid-IR.  Both \\citet{hin02} and \\citet{ow03} examined the mid-IR colors of small samples of stars with LSPs in the solar vicinity using IRAS data.  They found no evidence for a mid-IR excess that would indicate the presence of unusual amounts of circumstellar dust.  With the completion of the Spitzer Space Telescope SAGE survey in the LMC \\citep{meix06,blum06}, it is now possible to search for the presence of a mid-IR excess in large samples of stars with LSPs, good contemporary light curves and a known distance.  We now describe such a search. ",
        "conclusions": "We have shown that luminous red giant stars that exhibit Long Secondary Periods have larger mid-IR fluxes than similar stars without LSPs.  This suggests that the LSP induces additional mass loss from the red giant, with consequent dust formation and an increase in the mid-IR flux.  The mid-IR flux excess is only weakly dependent on the amplitude of the light variation of the LSP.  A comparison of the near-IR $J$-$K$ color with the mid-IR $K$-$[24]$ and $K$-$[8]$ colors indicates that the dust is not in a spherically symmetric distribution, and is perhaps in a disk.  Although dust is associated with the LSP phenomenon, it is still not clear what causes the LSP in the first place.  In a binary model, where the velocity curves indicate predominantly eccentric orbits \\citep{nic09}, mass transfer from the red giant near periastron could lead to a circumbinary disk.  A non-radial pulsation mode of low degree could also produce non-spherical mass loss and a non-spherical dust distribution.  The semi-periodic dust ejection events predicted by \\citet{win94} and \\citet{hof95} produce, and indeed require, circumstellar dust.  However, this phenomenon has not been found to occur in theoretical models at the low luminosities where many of the LSPs are observed (even below the tip of the FGB).  Furthermore, it is also not clear why the photospheric velocity should vary with the LSP for this purely circumstellar process.  One additional characteristic of LSPVs is that they have a chromosphere which varies with the LSP \\citep{wok04}.  It is likely that both the chromosphere and the excess circumstellar dust are manifestations of the influence of the LSP phenomenon on matter above the stellar photosphere. Magnetic effects are a possible source of the chromosphere, but a search for magnetic fields in two solar-vicinity LSPVs has put an upper limit of 100 Gauss on magnetic fields covering more than about 10\\% of the surface of these stars \\citep{wmwn09}.  Currently there is no evidence that magnetic fields are the source of the chromosphere.  In summary, we have shown here that the LSP phenomenon produces excess circumstellar dust compared to stars without LSPs.  Unfortunately, since all the postulated models for the LSP are potentially capable of causing mass ejection, the present detection of excess dust around LSPVs does not help us distinguish between the various models."
    },
    "0910/0910.3530_arXiv.txt": {
        "abstract": "Based on a {new} estimation of their thickness, the global properties of relativistic slim accretion disks are investigated in this work. The resulting emergent spectra are calculated using the relativistic ray-tracing method, in which we neglect the self-irradiation of the accretion disk. The angular dependence of the disk luminosity, the effects of the heat advection and the disk thickness on the estimation of the black hole spin are discussed. Compare to the previous works, our improvements are that we use the self-consistent disk equations and we consider the disk self-shadowing effect. We find that at the moderate accretion rate, the radiation trapped in the outer region of the accretion disks will escape in the inner region of the accretion disk and contribute to the emergent spectra. At the high accretion rate, for the large inclination and large black hole spin, both the disk thickness and the heat advection have significant influence on the emergent spectra. Consequently, these effects will influence the measurement of the black hole spin based on the spectra fitting and influence the angular dependence of the luminosity. For the disks around Kerr black holes with $a=0.98$, if the disk inclination is greater than $60^\\circ$, and their luminosity is beyond 0.2 Eddington luminosity, the spectral model which is based on the relativistic standard accretion disk is no longer applicable for the spectra fitting. { We also confirm that the effect of the self-shadowing is significantly enhanced by the light-bending, which implies that the non-relativistic treatment of the self-shadowing is inaccurate. According to our results, the observed luminosity dependence of the measured spin suggests that the disk self-shadowing significantly shapes the spectra of GRS 1915+105, which might lead to the underestimation of the black hole spin for the high luminosity states.} ",
        "introduction": "The standard accretion disk model \\citep[SSD,][]{1973A&A....24..337S} and its relativistic generalization \\citep{1973blho.conf..343N} are widely used in modeling the spectral energy distribution (SED) of both active galactic nuclei (AGNs) and the galactic black hole binaries. Recently, based on the relativistic SSD, the spins of the black holes in several galactic black hole binaries are measured by fittings of their SED \\citep{2005ApJS..157..335L,2006ApJ...636L.113S,2006ApJ...652..518M,2006MNRAS.373.1004M}. One major advantage of standard accretion disk model is its simplicity, i.e., it assumes that the fluid undergoes the Keplerian motion and the energy generated by the viscous heating is radiated locally. Besides, in this model, the disk terminates at the Innermost Stable {Circular} Orbit (ISCO), and the matter inside ISCO is cold.  However, when the accretion rate is relatively high ($\\dot M>0.1 \\dot M_{\\rm edd}$), in the inner region of the accretion disks, several effects will make these assumptions invalid. One such effect is the heat advection. When the vertical diffusion time scale is larger than the accretion time scale, substantial amount of heat is trapped inside the accretion flow and not able to escape to the infinity \\citep{1978MNRAS.184...53B,1982ApJ...253..873B}. This will make the local flux different from that from the SSD model and affect the profile of the radial temperature. Also, the advected heat will provide an additional pressure, which makes the disk rotation no longer Keplerian. Besides, because the heat trapped inside the accretion flow, the region inside ISCO is not cold any more.  The self-consistent treatment of these effects needs to model the energy and momentum balance in the disk, one widely used method is based on the radiative MHD simulation. Although the numerical simulations have the advantage of treating the physical effects self--consistently, the huge amount of computation makes it difficult to give the qualitative fitting of the observational data. One complementary approach is to solve the energy and momentum equations of the accretion flow in terms of the ordinary differential equations (ODEs). Compared to the simulations, this method can effectively take into account the dynamical effects, while maintaining its conciseness. In this work, we use the ODE form disk equations \\citep{1981AcA....31..283P,1988ApJ...332..646A,1993ApJ...412..254C,1997MNRAS.286..681P,1998ApJ...498..313G}, and explore the solutions at both the medium and high accretion rate \\citep[where the disk is call the ``slim disk\";][]{1988ApJ...332..646A}. We use the relativistic disk equations in our work \\citep{1996ApJ...471..762A,1998MNRAS.297..739B,2003MNRAS.338.1013S} to investigate the effects of the black hole spin.  The physical properties of the black hole and the accretion disks are deduced mainly from their observed SED. It is well known that the structure and the SED of the slim disk differ significantly from that of SSD \\citep{1999ApJ...516..420W,1999ApJ...522..839W}. The disk SED is influenced by a handful of effects, such as the Doppler effect, the relativistic beaming, the gravitational redshift, the gravitational -bending, and the disk self-shadowing. As the accretion rate increases, the disk becomes geometrically thick, and the disk self-shadowing \\citep{2007MNRAS.378..841W} becomes significant. In this work, we employ a ray-tracing approach \\citep[e.g.][]{1997PASJ...49..159F,1998NewA....3..647C} to explore the final spectra at the different viewing angles.  Due to their disk-like geometry, the radiation from the accretion disks is far from the isotropy, and the emergent luminosity is thus angular dependent \\citep[e.g.][]{2006MNRAS.372.1208H}. For a slab of optically thick material without velocity, the disk luminosity simply varies as $L\\sim \\cos\\theta$ in the flat spacetime. In fact, because of its thickness and rotation, the angular dependence of the disk luminosity is expected to be different from that of the simple slab, further analysis is thus needed. In this paper, we calculate  the angular dependence of the disk luminosity under different inclinations and the black hole spin.   { The spin of GRS 1915+105 is debated in the literature \\citep{2006ApJ...652..518M,2006MNRAS.373.1004M}. By fitting the disk blackbody component, \\citet{2006ApJ...652..518M} found that the spin is near extreme, while \\citet{2006MNRAS.373.1004M} found the spin is moderate, with $a \\sim 0.8$. To show how the spin determination is influenced by the heat advection and disk self-shadowing, we re-simulate the observation carried out by \\citet{2006ApJ...652..518M} with our model and confirm the luminosity dependence of the measured spin as found by \\citet{2006ApJ...652..518M}. }  In this work, the black hole mass in our canonical model is taken to be $10 M_{\\odot}$. However, due to the physical similarities between the accretion disks around the stellar-mass black holes and those around the supermassive black holes, our results can be easily generalized to the supermassive black holes and predict the SED of AGNs under the simple scaling.  This paper organized as follows. In section 2, we describe briefly our disk model and the ray--tracing method. In Section\\ref{secst}-Section \\ref{secend}, we present our numerical results of the global disk structure. In Section \\ref{respec}, we present our results on the SED of the disk, focusing on the differences between the SED of the dynamical disk model and that of the SSD. The angular dependence of the total luminosity is given in Section \\ref{angulardep}. In Section \\ref{reimp}, implications for estimating the black hole spin are discussed.  In Section \\ref{disc}, various observational implications of the disk self-shadowing as well as the spin of GRS 1915+105 are discussed. Finally, in Section \\ref{recon} we summarize our conclusions. ",
        "conclusions": "\\label{recon} In this work, we study the structure of the relativistic slim accretion disks and their emergent spectra. Besides the effect of the photon trapping, we find that when the accretion rate is appropriate, the energy trapped in the outer regions of the accretion disks can be re-radiated in the inner regions of the accretion disks. Another main result of this work is that we check the validity of the relativistic SSD. We find that at low accretion rate, such as $\\dot M \\sim 1/4$ for $a=0$ and $\\dot M \\sim 1/16$ for $a=0.98$, the relativistic SSD is still valid to describe the disk structure, and the effect of heat advection will not significantly influence the global structure of the disk. At relatively high accretion rate such as $\\dot M = 2 \\dot M_{\\rm edd}$, these effects will significantly distort the emergent spectrum, and hence affect the measurement of the black hole spin. We also find that when the accretion rate is high, the effect of the disk self-shadowing plays an even more important role than that of the heat advection alone.   The emergent luminosity from the accretion disks is angular dependent, and the degree of anisotropy is dependent on the accretion rate. In this work, we investigate the angular dependence of the observed luminosity. Due to the physical similarities between the accretion disks of different black hole mass, our result can be applied to estimate the radiation anisotropy of AGNs.  To check in principle the effect differences of spectra on estimation of spin, we create ``fake'' data sets with the simulated data, and fit them with the KERRBB model using XSPEC. At the relatively low accretion rate, our fitting results suggest that the simple Keplerian model is reliable. At the relatively high accretion rate, our fitting results suggest that self-shadowing effect will lead to significant under-estimation of the black hole spin. {By making our calculation relevant to the case of GRS 1915+105, without introducing additional parameters, we are able to reproduce the decrease of measured spin with the increasing luminosity, which is found in \\citet{2006ApJ...652..518M}. This finding suggests that disk self-shadowing significantly shapes the spectra of GRS 1915+105.}   { There are several effects which we do not include in our calculations. It was claimed that the disk may have no time to relax to the thermodynamic equilibrium \\citep{2002ApJ...574..315O}, this effect will make our estimation of emergent flux (Equation (\\ref{rad})) invalid and make the disk thickness estimation inaccurate. In our calculations of the disk SED, we use a simplified model of the local spectrum and neglect the effect of limb--darkening and returning irradiation. Assuming a constant color correction may make our calculated spectra inadequate to be compared to the observation \\citep{2008ApJ...683..389D}, so we focus on analyzing the physical effects we have introduced. The limb-darkening effect will influence the spectra and the angular dependence of the emergent luminosity, while the returning irradiation will have some effects on both the disk structure and the emergent spectra. As \\citet{2005ApJS..157..335L} have calculated, these two effects will contribute several percents to the final results. Despite of these shortcomings, our results and the trend of the spectral evolution are still robust when the inclination angle is not too large. }"
    },
    "0910/0910.3689_arXiv.txt": {
        "abstract": "At present none of galactic chemical evolution (GCE) models provides a self-consistent description of observed trends for all iron-peak elements with metallicity simultaneously. The question is whether the discrepancy is due to deficiencies of GCE models, such as stellar yields, or due to erroneous spectroscopically-determined abundances of these elements in metal-poor stars. The present work aims at a critical reevaluation of the abundance trends for several odd and even-$Z$ Fe-peak elements, which are important for understanding explosive nucleosynthesis in supernovae. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.4583_arXiv.txt": {
        "abstract": "The \\emph{Fermi Gamma-Ray Space Telescope} reveals a diffuse inverse Compton signal in the inner Galaxy with a similar spatial morphology to the microwave haze observed by WMAP, supporting the synchrotron interpretation of the microwave signal.  Using spatial templates, we regress out $\\pi^0$ gammas, as well as IC and bremsstrahlung components associated with known soft-synchrotron counterparts.  We find a significant gamma-ray excess towards the Galactic center with a spectrum that is significantly harder than other sky components and is most consistent with IC from a hard population of electrons. The morphology and spectrum are consistent with it being the IC counterpart to the electrons which generate the microwave haze seen at WMAP frequencies.  In addition, the implied electron spectrum is hard; electrons accelerated in supernova shocks in the disk which then diffuse a few kpc to the haze region would have a softer spectrum. We describe the full sky \\emph{Fermi} maps used in this analysis and make them available for download. ",
        "introduction": "The most detailed and sensitive maps of diffuse microwave emission in our Galaxy have been produced by the \\emph{Wilkinson Microwave Anisotropy Probe} (WMAP).  An analysis of the different emission mechanisms in these maps uncovered a microwave ``haze'' towards the Galactic center (GC) that has roughly spherical morphology and radius $\\sim4$ kpc \\citep{Finkbeiner:2003im}.  Since its discovery, \\cite{Finkbeiner:2004us} and \\cite{Dobler:2008ww} have argued that the microwave haze is hard synchrotron emission due to the fact that alternative hypotheses such as free-free (thermal bremsstrahlung of the ionized gas), spinning dust, or thermal dust have difficulty explaining its morphology, spectrum, or both.  However, the 23-33 GHz spectrum of the haze synchrotron is harder than that expected from diffused electrons originally accelerated by supernova shocks in the plane \\citep{Dobler:2008ww}.  Some variation is expected in the synchrotron spectrum, but generally in the sense that it should be \\emph{softer} at higher frequencies since the electrons lose energy preferentially at high energies as they diffuse from their source. While there are significant uncertainties in the spectrum of the haze \\citep[see the discussion in][]{Dobler:2008ww}, the data require that the diffused spectrum be roughly as hard as the expected \\emph{injection} spectrum from first-order Fermi acceleration at supernova shock fronts (number density $dN/dE \\propto E^{-2}$).  That is, if the electrons were produced in shocks in the disk, then they would have to have undergone no diffusive energy losses over a $\\sim4\\pi/3 \\ (4$ kpc$)^3$ volume, which seems unlikely.  Furthermore, there is significant emission in WMAP 23, 33, and 41 GHz bands from electrons that \\emph{were} generated in SN shocks; this emission has a very disk-like morphology (and softer spectral index which is consistent with shock acceleration), while the haze has a more spherical morphology.\\footnote{Hereafter ``spherical'' is taken to mean ``not disk-like'' -- if anything, the haze is non-spherical in the direction perpendicular to the disk (see \\refsec{residualmaps}).}  Together, the haze spectrum and morphology imply either (1) a new class of objects distributed in the Galactic bulge and largely missing from the disk; (2) significant acceleration from shocks several kpc off the plane towards the Galactic center; or, perhaps most intriguingly, (3) a new electron component from a new physical mechanism. The claim that novel physics or astrophysics is required to explain the WMAP data is called the \\emph{haze hypothesis} to distinguish it from the two null hypotheses: (1) that the microwave haze is not synchrotron, but rather some combination of free-free and spinning dust; and (2) that the haze is synchrotron, but the electron spectrum required is not unusual.  In this work we do not address the origin of the electrons, but instead consider what the data from the \\emph{Fermi Gamma-ray Space Telescope}\\footnote{See \\texttt{http://fermi.gsfc.nasa.gov/ssc/data/}} imply for their existence.  \\bpm \\includegraphics[width=0.49\\textwidth]{plots/skymap-002.0-005.0GeV.eps} \\includegraphics[width=0.49\\textwidth]{plots/skymap-005.0-010.0GeV.eps}  \\includegraphics[width=0.49\\textwidth]{plots/skymap-010.0-020.0GeV.eps} \\includegraphics[width=0.49\\textwidth]{plots/skymap-020.0-050.0GeV.eps}  \\includegraphics[width=0.49\\textwidth]{plots/skymap-050.0-100.0GeV.eps} \\includegraphics[width=0.49\\textwidth]{plots/expmap-002.0-005.0GeV.eps} \\caption{ \\emph{Top and bottom left:} Full sky \\emph{Fermi} $\\gamma$-ray maps in various energy bins.  The mask includes the 3-month \\emph{Fermi} point source catalog as well as the LMC, SMC, Orion-Barnard's Loop, and Cen A.  \\emph{Bottom right:} Exposure time map with our mask overlaid and stretched to $0-100\\%$ of the peak exposure time.  The variation in the exposure is a small modulation -- even setting it to unity does not change our qualitative results. } \\label{fig:fermimaps} \\epm  Electron cosmic rays at $10-100$ GeV primarily lose energy in the diffuse interstellar medium by producing synchrotron microwaves and inverse-Compton (IC) scattered gammas.  Synchrotron losses are proportional to magnetic field energy density, $U_B = B^2/8\\pi$, while IC losses are proportional to the interstellar radiation field energy density, $U_{ISRF}$, in the Thomson limit, and less in the Klein-Nishina limit.  Bremsstrahlung off the ambient gas also occurs, but is expected to be sub-dominant in the regions of interest. Therefore, the best test of the haze hypothesis is to search for IC gammas in the \\emph{Fermi} data, which was studied in the context of dark matter signals by \\cite{Cholis:2008wq,Zhang:2008tb,Borriello:2009fa,Cirelli:2009vg,Regis:2009md,Belikov:2009cx,Meade:2009iu}.  Previous studies of high latitude gamma-ray emission have reported measurements of an excess of emission above the Galactic plane, most notably SAS-2 \\citep{Fichtel:1975fi,Kniffen:1981kn}, COS-B \\citep{Strong:1984cb}, and EGRET \\citep{Smialkowski:1997sm,Dixon:1998di}.  However, in each case the experiments either covered an insufficient energy range (SAS-2 and COS-B) or did not have sufficient sensitivity and angular resolution (EGRET) to permit a spatial correlation with the WMAP haze. \\emph{Fermi} overcomes both of these obstacles and allows us to search for the gamma-ray counterpart to the microwave haze for the first time.  The presence of an IC signal at the expected level can confirm that the microwave haze is indeed synchrotron, ruling out the first null hypothesis.  From the spectrum of the IC, we can estimate the electron spectrum required to make the signal, addressing the second null hypothesis.  For example, the presence of $\\sim 50~\\GeV$ IC photons requires electrons of $E>50~\\GeV$, perhaps much greater. Furthermore, IC photons provide valuable information about the spatial distribution (disk vs.\\ bulge) of the source of these particles, which in turn can constrain hypotheses about their origin.\\footnote{The synchrotron haze depends on the Galactic magnetic field while the IC haze depends on the Galactic ISRF, and so the two morphologies should not be identical, and could be quite different.}  In \\S 2 we briefly review the \\emph{Fermi} data, describe our map-making procedure, and display full-sky maps at various energy ranges.  In \\S 3 these maps are analyzed with template correlation techniques and resultant residual maps and spectra are shown. Finally, \\S 4 presents our interpretation of the signals and discusses potential sources of contamination.  Estimates for the possible contamination from unresolved point sources are given in \\refapp{pointsources}. Appendix \\ref{sec:datarelease} details the creation and processing of the gamma-ray sky maps used in this analysis, and provides instructions for downloading them. Appendix \\ref{sec:likelihood} contains a discussion of Poisson likelihood analysis on smoothed maps. ",
        "conclusions": "We have presented full sky maps generated from photon events in the first year data release of the \\emph{Fermi Gamma-Ray Space Telescope} (see Appendix \\ref{sec:datarelease} for data processing details).  Using a template fitting technique, we have approximated both the spectrum and morphology of the well known gamma-ray emission components at \\emph{Fermi} energies.  The SFD dust map was used to trace the $\\pi^0$ decay gammas generated by collisions of cosmic ray protons with the ISM, while the Haslam 408 MHz map was used to trace inverse Compton (IC) scattered photons from interactions of supernova shock-accelerated electrons ($\\sim 1-10$ GeV) with the interstellar radiation field (ISRF).  Bremsstrahlung radiation, generated by interactions of these electrons with the ISM, should be approximately traced by some combination of these two maps.  Although our template fitting technique is subject to significant uncertainties due to uncertain line of sight gas and CR distributions, a robust positive residual has been identified.  This excess diffuse emission is centered on the Galactic center, and can be parameterized by a simple two-dimensional Gaussian template $(\\sigma_\\ell=15\\degree, \\sigma_b=25\\degree)$. The template-correlated spectrum of this emission is significantly \\emph{harder} than either $\\pi^0$ emission or IC from softer electrons, whose fitted spectra agree well with models.  This harder spectrum coupled with the distinct spatial morphology of the gamma-ray and microwave haze are evidence that these electrons originate from a \\emph{separate component} than the softer SN shock-accelerated electrons.  The gamma-ray excess is almost certainly the IC counterpart of the microwave haze excess described by \\cite{Finkbeiner:2003im} and \\cite{Dobler:2008ww}.  Although it is still possible that a significant fraction is prompt photons from WIMP annihilations \\cite[e.g. the 200 GeV wino advocated by][]{Grajek:2008pg} such explanations are difficult to reconcile with the spatial similarity to the WMAP haze (see \\reffig{fermi_comp_wmap}).  The simplest hypothesis is that the signal is mainly IC from the same electrons that produce the WMAP haze synchrotron.  This addresses the stubborn question about the origin of the WMAP haze. Until recently, it has been argued that the WMAP haze had alternative explanations, such as free--free emission from hot gas or spinning dipole emission from rapidly rotating dust grains.  However, the existence of this IC signal proves that the microwave haze is indeed synchrotron emission from a hard electron spectrum.  \\emph{Fermi} LAT photon data are contaminated by particle events, especially at high energies.  We have taken care to account for the isotropic background resulting from extragalactic sources, cosmic ray contamination, and heavy nuclei contamination and found that this background, though significant, is below the observed IC excess even up to 100 GeV.  \\emph{Particle contamination is extremely unlikely to mimic the observed signal.}  The LAT collaboration continues to refine the cuts used to define ``diffuse class'' events, and plans to release a cleaner class of events in coming months.  This, along with a new public version of GALPROP, including updated ISRF models, will allow a more sophisticated analysis than that presented in this paper.  We eagerly await the release of these software and data products.  The spectrum and morphology of both the microwave and gamma-ray haze constrain explanations for the source of these electrons.  There have been speculations that the microwave haze could indicate new physics, such as the decay or annihilation of dark matter, or new astrophysics, such as a GRB explosion, an AGN jet, or a spheroidal population of pulsars emitting hard electrons. We do not speculate in this paper on the origin of the haze electrons, other than to make the general observation that the roughly spherical morphology of the haze makes it difficult to explain with any population of disk objects, such as pulsars.  The search for new physics -- or an improved understanding of conventional astrophysics -- will be the topic of future work.  \\vskip 0.15in {\\bf \\noindent Acknowledgments:} We acknowledge helpful conversations with Elliott Bloom, Jean-Marc Casandjian, Carlos Frenk, Isabel Grenier, Igor Moskalenko, Simona Murgia, Troy Porter, Andy Strong, Kent Wood. This work was partially supported by the Director, Office of Science, of the U.S.  Department of Energy under Contract No. DE-AC02-05CH11231. NW is supported by NSF CAREER grant PHY-0449818, and IC and NW are supported by DOE OJI grant \\# DE-FG02-06ER41417. IC is supported by the Mark Leslie Graduate Assistantship.  DF and TS are partially supported by NASA grant NNX10AD85G. TS is partially supported by a Sir Keith Murdoch Fellowship from the American Australian Association.  This research made use of the NASA Astrophysics Data System (ADS) and the IDL Astronomy User's Library at Goddard.\\footnote{Available at \\texttt{http://idlastro.gsfc.nasa.gov}}   \\clearpage"
    },
    "0910/0910.1190_arXiv.txt": {
        "abstract": "The amount of available telescope time is one of the most important requirements when selecting astronomical sites, as it affects the performance of ground-based telescopes. We present a quantitative survey of clouds coverage at La Palma and Mt.Graham using both ground- and satellite-based data. The aim of this work is to derive clear nights for the satellite infrared channels and to verify the results using ground-based observations. At La Palma we found a mean percentage of clear nights of $62.6$ per cent from ground-based data and $71.9$ per cent from satellite-based data. Taking into account the fraction of common nights we found a concordance of $80.7$ per cent of clear nights for ground- and satellite-based data. At Mt.Graham, we found a $97$ per cent agreement between Columbine heliograph and night-time observing log. From Columbine heliograph and the {\\it Total Ozone Mapping Spectrometer-Ozone Monitoring Instrument (TOMS-OMI)} satellite, we found that about $45$ per cent of nights were clear, while satellite data (GOES, TOMS) are much more dispersed than those of La Palma. Setting a statistical threshold, we retried a comparable seasonal trend between heliograph and satellite. ",
        "introduction": "The identification and characterization of a site for the future European Large Telescope (E-ELT) is a key issue. Moreover a quantitative survey of cloud cover for the areas selected as candidate sites for the telescope is and will continue to be an essential part of the process of site selection for future large telescopes in the same class as the E-ELT.  In fact, the performance of large telescopes at optical and infra-red wavelengths is critically dependent on atmospheric cloud cover. Cloud cover is a key parameter at the time of site selection and also affects scientific output during the life of the telescope. For instance, a night-time seasonal trend of fewer clear sky can reduce regular access to the sky.  Typically it is possible to quantify the presence of clouds at telescope sites using ground-based observations that provide a real time knowledge of the atmospheric condition. The fraction of clear sky can be determined using either instruments (i.e all-sky cameras) or observer estimates. Long-term records from many ground-based telescopes, which list the number of nights available for observing, are now accessible and it is possible to begin a reliable statistical study. However, this technique alone is not suitable for identifying future candidate sites where there are no telescopes.  An easy evolution of this analysis is the use of meteorological satellites that provide measurements of cloud cover and other critical parameters for site testing covering large areas with different spatial and temporal resolution. Taking into account that, most of the meteorological satellites  are equipped with similar instrumentation, it is not difficult to compare distinct sites observed by two or more different satellites. Additionally, since satellite data archives now cover long time periods, it is possible to have for each site the trend of these parameters in both long and short time scale. Erasmus and Sarazin (\\cite {erasmus02}) were among the first to demonstrate the successful application of satellite data for monitoring, comparison and forecasting evaluations. Erasmus and van Rooyen (\\cite {erasmus06}) quantified cloud cover at La Palma using Meteosat satellite and validating them using the ground based measurements taken at Carlsberg Meridian Telescope (CMT).  In this paper we present the results of a study of cloud cover using satellite- and ground-based data obtained in two different important astronomical sites: the Observatorio del Roque de Los Muchachos (ORM) located in La Palma (in Canary Islands), hosting several international telescopes and among of them the Galileo National Telescope (TNG), and Mount Graham (in Arizona), hosting the Large Binocular Telescope (LBT). The results are compared with Erasmus and van Rooyen (\\cite {erasmus02}) and Erasmus and Sarazin (\\cite {erasmus06}). The present paper is organized as follows. In Section 2, we describe both the ground and satellite data bases adopted. In Section 3, we describe the satellite data acquisition procedure, and in Section 4 we show the data reduction procedure. Section 5 gives a discussion of the results. ",
        "conclusions": " \\begin{itemize} \\item \\textbf It is possible to define a threshold in satellite emissivity to select clear nights with an uncertainty of $20$ per cent. \\item \\textbf A good correlation exists between GOES 12 satellite and ground-based data. \\item \\textbf Using the common selected nights we found that $80.7$ per cent are classified clear from both GOES 12 satellite and ground logbook. \\item \\textbf The marginal increase to $81.1$ per cent of the concordance obtained including calima events confirm that the satellite is able to distinguish only the presence of clouds. \\end{itemize} A further analysis of the satellite data (e.g.a wider field, the simultaneous use of different channels, etc.) is suggest to improve the prediction of clear nights. Furthermore, we are studying the fraction of clear nights lost for high humidity or strong wind.     \\subsection"
    },
    "0910/0910.1159_arXiv.txt": {
        "abstract": " ",
        "introduction": "One of the complications in VLBI, over connected array interferometers, arises from the completely unrelated atmospheric conditions that the wavefronts propagate through before reaching the widely separated antennae. Self calibration procedures, which are the standard VLBI analysis technique for imaging radio sources, rely on closure relations to remove the station dependent complex gain factors that characterize the phase errors at each antenna.  A direct detection of the source with good signal to noise ratio is required within every segment of the coherence integration time-interval. This time interval is set by the stability of the instrument and, dominantly, the atmospheric turbulence.  An important consequence of the use of phase closure is that information on the absolute position of the source is lost, preventing the measurement of astrometric quantities.\\\\  The application of phase-referencing  techniques, to the analysis of interleaving observations of the program source and a nearby calibrator, preserves the information on the angular separation on the sky and provides high precision relative-astrometry (Alef 1988, Beasley \\& Conway, 1995). At the observations, the scans on the scientifically interesting source, the target, are interleaved (within the coherence integration time) with observations of a calibrator source, the reference. The antenna-based corrections derived from the self-calibration analysis of the reference source observations are transferred for the calibration of the target source. Next, the target dataset is Fourier transformed without any further calibration to yield a phase referenced map of the target source, where the position offset of the peak from the center provides a precise measurement of the relative separation between both sources. The propagation of astrometric errors in the phase referencing analysis is strongly dependent on the angular separation between the target and reference sources, and range between the micro-arcsec and tenths of milli-arcsec accuracy. This phase referencing technique, from now on referred as ``conventional phase referencing'', is well established and has been used to provide high precision astrometric measurements of (relative) source positions in cm-VLBI observations.  It would be highly desirable to extend this capability to the mm-VLBI regime, yet at the highest frequencies the observations are sensitivity limited: the instruments are less efficient, the sources are intrinsically weaker, and the phase coherence integration times are severely constrained by the rapid atmospheric phase fluctuations due to the variations (spatial and temporal) of the water vapor content in the troposphere. In particular, the coherence time is too short to allow an antenna to switch its pointing direction between pairs of sources, in all but the most exceptional cases (Porcas \\& Rioja, 2002), within that time range. The lack of suitable reference sources in mm-VLBI makes it almost impossible to apply ``conventional phase referencing'' techniques in the high frequency (i.e. significantly above 43\\,GHz) domain. \\\\  Therefore it would be hugely beneficial if the calibration could be performed at a lower, easier, frequency and used for data collected at a higher frequency. That is, to transfer the calibration terms (for phase/delay/rate VLBI observables) derived at a different frequency rather than at a different source as in ``conventional phase-referencing''. It should be noted that frequency switching can be performed much faster than source switching at the VLBA.  Moreover the duty-cycle is now determined by the coherence time at the lower frequency. The low frequency phases provide `connection' but of course can not correct for variations faster than the duty cycle. This requires co-temporal dual frequency observations as provided by the next generation of VLBI antennae and arrays which are able to co-observe at different frequency bands, e.g. the Yebes 40m antenna and the Korean VLBI Network.\\\\     The feasibility of multi-frequency observations to correct the non-dispersive tropospheric phase fluctuations in the high frequency regime has been studied for some time.  It relies on the fact that such fluctuations will be linearly proportional to the observing frequency, and hence it should be possible to use a scaled version of the calibration terms derived from the analysis of observations at a lower frequency (where more and stronger sources are available, with longer coherence integration times and better antenna performance), to calibrate higher frequency observations. It is a kind of phase referencing, between observations at two frequencies, that we call ``frequency phase transfer'' ({\\sc FPT}). Among the earliest references we found are ``Phase compensation experiments with the paired antennas method 2. Millimeter-wave fringe correction using centimeter-wave reference'' (Asaki, \\etal\\ 1998) with the Nobeyama millimeter array (NMA), and ``Tropospheric Phase Calibration in Millimeter Interferometry'' (Carilli \\& Holdaway, 1999) for application with the Very Large Array (VLA). In ``VLBI observations of weak sources using fast frequency switching'', Middelberg \\etal\\ (2005) applied this frequency phase transfer technique to mm-VLBI observations. They achieved a significant increase in coherence time, resulting from the compensation of the rapid tropospheric fluctuations, but failed to recover the astrometry, due to the remaining residual dispersive terms.  Our proposed {\\sc Source/Frequency phase referencing} method endows this approach with astrometric capability for measuring frequency dependent source positions (``core-shifts'') by adding a strategy to estimate the ionospheric (and other) contributions. In ``Measurement of core-shifts with astrometric multi-frequency calibration'' (Rioja \\etal\\ 2005) we applied it to measure the ``core-shift'' of quasar 1038+528 A between S and X-bands (8.3/2.2 GHz), and validate the results by comparison with those from standard phase referencing techniques at cm-VLBI, where both methods are equivalent. \\\\ Here we present a demonstration of successful application of the {\\sc sfpr} method to astrometric mm-VLBI, a much more challenging frequency regime where conventional phase referencing fails.  Also, the basis of the method, and details on the scheduling and data analysis are described. \\\\ This method opens a new horizon with targets and fields suitable for high precision astrometric studies with VLBI, especially at high frequencies where severe limitations imposed by the rapid fluctuations in the troposphere prevent the use of conventional phase referencing techniques. In addition this method can be applied to Space VLBI, where accurate orbit determination is a significant issue. This method results in perfect correction of frequency independent errors, such as those arising from the uncertainty in the reconstruction of the satellite orbit. The application to the space mm-VLBI mission VSOP-2 is described in detail in Rioja \\& Dodson (2009). ",
        "conclusions": "We have proposed and demonstrated a new method of astrometric VLBI calibration, suitable for mm-VLBI. It uses dual frequency observations to removed non-dispersive contributions. The additional step required to remove the ionospheric, and all other slowly varying dispersive terms, is done by including another source to cross calibrate with. Because the ionospheric patch size is very large at mm wavelengths one can use calibrators that lie a considerable distance from the source.  We have presented the results from two pairs of sources, one only $14^\\prime$ apart and the other 10$^o$ apart. A single pair is sufficient to demonstrate the method, however the astrometric solution (the offset from expected position) contains the contribution from both sources (as happens in standard phase-referencing as well). This problem, however, fulfills the closure condition, so three or more sources can be used to form a closure triangle and separate the contributions from each individual source."
    },
    "0910/0910.5066_arXiv.txt": {
        "abstract": "The form of the nuclear symmetry energy $E_s$ around saturation point  density leads to a different crust-core transition point in the neutron star and affect the crust properties. We show that the knowledge about $E_s$ close to the saturation point is not sufficient, because the very low density behaviour is relevant. We also claim that crust properties are strongly influenced by the  very high density behaviour of $E_s$, so in order to conclude about the form of low density part of the symmetry energy one must isolate properly  the high density part. ",
        "introduction": "One of the most intriguing quantity in the description of nuclear matter in neutron stars is the symmetry energy $E_s$, which is defined as follows \\beq E_{nuc}(n,x)=V(n) + E_s(n)\\, \\alpha^2 + {\\cal O}(\\alpha^4) \\label{Enuc} \\eeq where  $E_{nuc}(n,\\alpha)$ represents the energy of nucleonic matter per baryon as a function of baryon number density $n$ and the isospin asymmetry $\\alpha$, where $\\alpha\\!=\\!(n_n\\!-\\!n_p)/n$ and $n_n,n_p$ are the neutron and proton densities. At the  saturation point density $n_0=0.16 \\fm3$ the value of symmetry energy corresponds to the $a_4$ parameter in the Bethe-Weizs\\\"acker mass formula, and takes the value $E_s(n_0) = 30\\pm 1\\MeV$. Isoscalar part  of interactions is represented by the isoscalar potential $V(n)$ which is mainly responsible for the stiffness of the Equation of State (EoS).  Density dependence of $E_s$ is however highly uncertain both below and above saturation point $n_0$. This dependence is one of the goals of the experimental investigations carried on the radioactive beam colliders \\cite{Baran:2004ih,Li:2008gp}. This kind of facilities allow for research of nuclear matter with large isospin asymmetry. The analysis shown in  \\cite{Chen:2004si,Chen:2005ti,Chen:2007qb} put some constraints on the slope and curvature of $E_s(n)$ around $n_0$ however we are still far from the final conclusion about the global shape of the symmetry energy.  The role played by the $E_s$ in the context of various neutron star observables was emphasized in \\cite{Lattimer:2000nx}   and more detailed analysis was made in \\cite{Steiner:2004fi}. One of the first approach to the crust-core transition  was presented in  \\cite{Baym:1971ax} where the stability considerations were performed. This kind of analysis with an improved nuclear model was later used  in  \\cite{Pethick:1994ge}. In this work  authors suggested the different critical density in different nuclear models comes from their discrepancy in the neutron matter description, e.g. in the symmetry energy form. The direct connection of $E_s$ to the crust-core transition point was shown in \\cite{Kubis:2006kb} and possible phase separation in the inner core was analysed in \\cite{Kubis:2007zz}. In this work we would like to go along this line to emphasize that the very low density behaviour of the symmetry energy is essential for the crust-core transition point. It is especially interesting taking into account the recent experimental results \\cite{Kowalski:2006ju,Natowitz:2010ti} which show the symmetry energy still takes large values at densities very much below $n_0$. This result is in contrast to common conviction coming from various theoretical approaches that states the $E_s$ goes almost linearly to zero for low densities. So, the consequences of the symmetry energy with such unusual feature seems to be worth to be seen. ",
        "conclusions": ""
    },
    "0910/0910.2193_arXiv.txt": {
        "abstract": "We present an analysis of the luminosity distances of Type Ia Supernovae from the Sloan Digital Sky Survey-II (SDSS-II) Supernova Survey in conjunction with other intermediate redshift ($z<0.4$) cosmological measurements including redshift-space distortions from the Two-degree Field Galaxy Redshift Survey (2dFGRS), the Integrated Sachs-Wolfe (ISW) effect seen by the SDSS, and the latest Baryon Acoustic Oscillation (BAO) distance scale from both the SDSS and 2dFGRS. We have analysed the SDSS-II SN data alone using a variety of ``model-independent\" methods and find evidence for an accelerating universe at  $>$97\\% level from this single dataset. We find good agreement between the supernova and BAO distance measurements, both consistent with a $\\Lambda$--dominated CDM cosmology, as demonstrated through an analysis  of the distance duality relationship between the luminosity ($d_L$) and angular diameter ($d_A$) distance measures. We then use these data to estimate $w$ within this restricted redshift range ($z<0.4$). Our most stringent result comes from the combination of all our intermediate--redshift data (SDSS-II SNe, BAO, ISW and redshift--space distortions), giving $w = -0.81^{+0.16}_{-0.18}(\\mathrm{stat}){\\pm0.15}(\\mathrm{sys})$ and $\\Omega_M=0.22^{+0.09}_{-0.08}$ assuming a flat universe. This value of $w$, and associated errors, only change slightly if curvature is allowed to vary, consistent with constraints from the Cosmic Microwave Background. We also consider more limited combinations of the geometrical (SN, BAO) and dynamical (ISW, redshift-space distortions) probes. ",
        "introduction": "It is now widely believed that the late-time expansion of the universe is accelerating.  General Relativity (GR) implies that the acceleration is driven by ``dark energy\" (DE) --- an unknown energy component in the universe with a negative effective pressure.  Describing dark energy by an equation-of-state parameter of $w(z) = p/(\\rho c^2 )$ requires that $w<-1/3$.  Alternatively, accelerated expansion could be an indication that GR is not the correct theory of gravity or that we have applied GR incorrectly in a cosmological context (see recent reviews of DE by \\citealt{2003RvMP...75..559P}, \\citealt{2006astro.ph..5313U}, \\citealt{2007AIPC..957...21C}, and \\citealt{2008arXiv0803.0982F}).  Over the last decade, the most direct way of studying this acceleration of the expansion of the universe, and therefore DE, has been using Type Ia Supernovae (SNe), as they have been shown by many authors to be well-calibrated ``standard candles\" in the universe, i.e., their relative distances can be determined from the dependence of their peak luminosity on the shape of the light curve. This method was used to great effect by astronomers in 1998 to provide the first evidence for an accelerated universe (\\citealt{1998AJ....116.1009R, 1999ApJ...517..565P}; see \\citealt{2005ASSL..332...97F} for a review).  Briefly, a Type Ia supernova occurs when a white dwarf star in a close binary system accretes enough mass from its companion to undergo a thermonuclear explosion in the core. Both \\citet{1993ApJ...413L.105P} and \\citet{1993PASP..105..787H} have shown that such explosions can serve as consistent light sources in the universe to high accuracy. This is achieved  by transforming the measured light curve of the explosion into the rest frame of the supernova (so called K-corrections) and then correcting the luminosity at maximum as a function of the shape of the rest-frame light curve.  Several techniques now exist for fitting SN light curves known under different acronyms ($\\Delta m_{15}$, \\citealt{1996AJ....112.2391H}; MLCS, \\citealt{1996ApJ...473...88R}; stretch, \\citealt{1997ApJ...483..565P}; CMAGIC, \\citealt{2003ApJ...590..944W}; BATM, \\citealt{2003ApJ...594....1T}; SALT, \\citealt{2005A&A...443..781G}; $\\Delta C_{12}$, \\citealt{2006ApJ...645..488W}; SALT2, \\citealt{2007A&A...466...11G}; SiFTO, \\citealt{2008ApJ...681..482C}). In this analysis, we consider MLCS2k2 \\citep{2007ApJ...659..122J, 1996ApJ...473...88R}, which is among the most commonly used, and best tested, of these various techniques.  Recently, several dedicated SN surveys have been carried out to confirm and extend the earlier detections of an accelerating universe (HST, \\citealt{2004ApJ...607..665R,2007ApJ...659...98R}; SNLS, \\citealt{2006A&A...447...31A}; ESSENCE, \\citealt{2007ApJ...666..694W}) as well as new compilations of existing SN datasets \\citep{2007ApJ...666..716D, 2008arXiv0804.4142K, 2009arXiv0901.4804H}. In addition to supernovae, observations of Baryon Acoustic Oscillations (BAO) can be used to measure distances in the universe \\citep{2003ApJ...594..665B, 2003ApJ...598..720S, 2003PhRvD..68f3004H}. The BAO are caused by sound waves in the early universe which leave a preferred scale in the distribution of matter equal to the sound horizon at recombination. Today, this scale corresponds to $\\sim100/h$ Mpc (Hubble constant at present: $H_0 = 100h$~km/s/Mpc) and can thus be used as a ``standard ruler\" throughout the universe. The BAO signature has been detected in the clustering of galaxy clusters by \\citet{2001ApJ...555...68M}, in the correlation of galaxies in the Sloan Digital Sky Survey (SDSS, \\citealt{2000AJ....120.1579Y}) by \\citet{2005ApJ...633..560E}, \\citet{2006AA...459..375H}, \\citet{2007MNRAS.378..852P}, and \\citet{2007MNRAS.374.1527B}, and in the Two-degree Field Galaxy Redshift Survey (2dFGRS, \\citealt{2001MNRAS.328.1039C}) by \\citet{2005MNRAS.362..505C}.  In addition to the geometrical methods discussed above, observations of the dynamical properties of matter can provide constraints on the matter density of the universe and $w$ (assuming General Relativity is the appropriate theory of gravity). For example, the growth rate of structure in the universe can be observed via the coherent infall of galaxies into large clusters and superclusters of galaxies seen in redshift surveys \\citep{1987MNRAS.227....1K}. Also, the growth of structure can be measured using the late-time Integrated Sachs-Wolfe (ISW) effect \\citep{1967ApJ...147...73S}, which has now been detected to high significance through the cross-correlation of galaxy surveys with the Cosmic Microwave Background (CMB) (see \\citealt{2008PhRvD..77l3520G} and references therein). The ISW is sensitive to deviations from a matter-dominated, Einstein-de Sitter universe ($\\Omega_M=1$, where $\\Omega_M$ is the matter density at present divided by the critical density.).  Taken together, the present combination of cosmological measurements suggests we live in a flat univers, dominated by a cosmological constant ($\\Lambda$) with the energy density in matter and $\\Lambda$ known to a statistical accuracy of better than a few percent \\citep[see][]{2009ApJS..180..306D}. However, several of these cosmological measurements, especially SNe, are becoming limited by their systematic uncertainties which are now dominating, e.g., \\citet{2009arXiv0901.4804H} showed that the best combination of available SNe and BAO measurements provide $1+w=0.013^{+0.066}_{-0.068}$ but with a systematic uncertainty of $0.11$. Therefore, it is clear that future cosmological surveys must resolve these systematic errors through new observations and better analysis methods to mitigate their effect.  This paper is one of three complementary papers focused on the cosmological analysis of a new sample of  intermediate supernova distances recently obtained by the SDSS-II Supernova Survey (see Section \\ref{chapter_1st_year_data} for details). Our analysis differs from those presented in our companion papers of \\citet{kessler} and \\citet{davis2008}, as we first study the cosmological information obtained solely from the SDSS-II SN sample, and then in combination with other cosmological probes over the same redshift range ($z<0.4$).  Alternatively,  \\citet{kessler} presents a detailed examination of the impact of both statistical and systematic errors on deriving standard cosmological constraints based on the combination of the SDSS-II SN with most of the currently available high and low redshift SNe and which are all analysed in a consistent way. \\citet{davis2008} then use the same compilation of data to study an expanded set of exotic cosmological models, in combination with a wider variety of other cosmological information. Our approach is also complementary to the many other analyses in the literature (e.g. \\citealt{2007ApJ...666..716D}, \\citealt{2008arXiv0804.4142K}, \\citealt{2009arXiv0901.4804H}) that have used data from all possible sources.  In our approach we concentrate on the information from cosmological measurements that cover the same range of redshifts as the SDSS SN sample. Our aim is not to derive the most stringent limit on cosmological parameters available, but rather to verify that if we restrict ourself to probes coming from a small and similar redshift slice the results on cosmological parameters remain stable and consistent. This approach is warranted because of the growing emphasis on controlling systematic uncertainties in the analysis of cosmological data. There are clearly a number of systematic uncertainties that could affect the use of SNe as cosmological probes which likely depend on, or change with, redshift, including SN evolution (e.g., changes in the metallicity of progenitor stars \\cite{2003ApJ...590L..83T, 2009ApJ...691..661H, 2009ApJ...693L..76S}), intergalactic dust \\citep{2007ApJ...664L..13C, 2008MNRAS.386..475H}, Malmquist bias and the effects of gravitational lensing and peculiar velocities \\citep{2006PhRvD..73l3526H}. Moreover, the photometric uncertainties associated with combining SN data from multiple surveys, over a range of redshifts, is already seen as the main limitation in using presently available datasets \\citep[see][]{2009arXiv0901.4804H}. Our analysis addresses this issue by focusing exclusively on the SDSS SN sample, which is derived from a well-understood and stable photometric system. The SDSS has a  relative photometric accuracy of better than $2\\%$ in griz, and  $4\\%$ for the u--band (Padmanabhan et al. 2007), while the absolute calibration is also known to be of the order of $1\\%$, leading to a homogeneous set of SN light-curves with high photometric accuracy \\citep[see][]{2008AJ....136.2306H}. This set of data is robust to uncertainties in light-curve fitting. For example, the MLCS2k2 and SALT2 fits to the SDSS-only sample are shown to agree well \\citep[see Section 10 in][]{kessler} which is not the case when the higher redshift SN samples are added.  The outline of this paper is as follows. In Section \\ref{chapter_1st_year_data}, we describe the SDSS-II SN data and use that data in Section 3 to study the cosmic acceleration in the Universe. Section 4 then compares the SDSS-II SN luminosity distances to the BAO distances from the SDSS and 2dFGRS, checking the distance duality relation. We then derive in Section 5 constraints on $w$ by combining the best-fit luminosity distances for SDSS-II SNe with the growth rate of structure measurements taken from \\citet{2003MNRAS.346...78H} and a new measurement of the ISW effect taken from  \\citet{2008PhRvD..77l3520G}. We conclude in Section \\ref{conclusions}.  \\begin{figure} \\center \\includegraphics[width=80mm] {fig1.eps} \\caption{Residual Hubble diagram with respect to an empty universe for the 103 Type Ia SNe from the first year of operation of the SDSS-II SN Survey. The red line shows a $\\Lambda$CDM model with $(\\Omega_M, \\Omega_\\Lambda) = (0.3, 0.7)$, similar to our best-fit model given in Table \\ref{table_w}. The blue line is the best fit to the data derived from the ``sliding window\" technique described in Section \\ref{sec:nonpa}, and the cyan and green shaded regions correspond to the $1\\sigma$ and $2\\sigma$ confidence intervals, respectively. The black line indicates an expansion history with a deceleration parameter of $q_0 = -0.34$ as described in Section \\ref{section_global_fit}. The lower panel shows the same parameterization but now converted to $D_V$ according to Eq. (\\ref{equation_dv}). The two data points represent the BAO measurements from \\citet{2009arXiv0907.1660P}.} \\label{fig_sdss_hubble} \\end{figure} ",
        "conclusions": "\\label{conclusions}  We present an analysis of the luminosity distances of Type Ia Supernovae from the Sloan Digital Sky Survey-II (SDSS-II) Supernova Survey in conjunction with other intermediate redshift ($z<0.4$) cosmological measurements including redshift-space distortions from the 2dFGRS, the ISW effect, and the BAO distance scale from both the SDSS and 2dFGRS. We have analyzed the SDSS-II SN luminosity distances using several 'model-independent' methods, including fitting the data using a $q(z)$ parameterization, principal components, and a non-parametric ``sliding window\" method. We find consistent results between all these methods that provides evidence for an accelerating universe based solely on the first-year SDSS-II SN data. The strongest evidence we find comes when we make the strongest assumptions, that $q_0$ is constant and the universe is flat which gives probability for acceleration of $>97\\%$.  We also compare our SDSS-II SN data with the local BAO measurements, and find they are in good agreement. This is in contrast with the findings of \\citet{2007MNRAS.381.1053P} which found tension between the two distance measures, but confirms the new BAO analysis of Percival et al. (2009) who note that this tension has now lessened. Taking this observation further, we test the distance duality relation, i.e., for any metric theory of gravity, we expect $d_L/(d_A (1+z)^2) = 1$. We see no evidence for a discrepancy from this relation (at the one sigma level) in contrast to previous claims for a potential violation on the $2\\sigma$ level as seen in \\citep{2004PhRvD..69j1305B, 2008JCAP...07..012L}. Finally, we present a new measurement of the equation-of-state parameter of dark energy using a combination of geometrical distances in the universe and estimates for the growth rate of structure. Our strongest constraint comes from the combination of all four data-sets discussed herein (SDSS-II SN, BAO, redshift-space distortions, ISW) with $w=-0.81^{+0.16}_{-0.18}(stat)$ and $\\Omega_M=0.22^{+0.09}_{-0.08}(stat)$ (assuming a flat universe).  However, the combination of just the SDSS-II SNe and the ISW measurements alone is almost as powerful in constraining these parameters (Table 1). Our results only change slightly if we allow curvature to vary, consistent with the CMB measurements (see Appendix \\ref{appendix_curvature}). We quote a systematic uncertainty of $\\Delta w =\\pm0.15$ based on the details of the MLCS2k2 light--curve fitter (see \\citealt{kessler} for a fuller discussion).  Thus we have shown that low-redshift cosmological probes give a self-consistent picture of the distance-redshift relation.  When combined with growth of structure and ISW at the same epoch that picture is consistent with $\\Lambda$CDM and re-enforces the complementarity amongst other data and analyses in the literature."
    },
    "0910/0910.5857_arXiv.txt": {
        "abstract": "We present a transport model that describes the orbital diffusion of asteroids in chaotic regions of the 3-D space of proper elements. Our goal is to use a simple random-walk model to study the evolution and derive accurate age estimates for dynamically complex asteroid families. To this purpose, we first compute local diffusion coefficients, which characterize chaotic diffusion in proper eccentricity ($e_p$) and inclination ($I_p$), in a selected phase-space region. Then, a Monte-Carlo-type code is constructed and used to track the evolution of random walkers (i.e.\\ asteroids), by coupling diffusion in ($e_p$,$I_p$) with a drift in proper semi-major axis ($a_p$) induced by the Yarkovsky/YORP thermal effects. We validate our model by applying it to the family of (490) Veritas, for which we recover previous estimates of its age ($\\sim 8.3$~Myr). Moreover, we show that the spreading of chaotic family members in proper elements space is well reproduced in our randomk-walk simulations. Finally, we apply our model to the family of (3556) Lixiaohua, which is much older than Veritas and thus much more affected by thermal forces. We find the age of the Lixiaohua family to be $155\\pm 36$~Myr. ",
        "introduction": "\\label{}  As first noted by \\citet{hirayama1918}, asteroids can form prominent groupings in the space of orbital elements. These groups, nowadays well-known as \\emph{asteroid families}, are believed to have resulted from catastrophic collisions among asteroids, which lead to the ejection of fragments into nearby heliocentric orbits, with relative velocities much lower than their orbital speeds. To date, several tens of families have been discovered across the whole asteroid Main Belt \\citep[e.g.][]{bend02,nes2006}. Also, families have been identified among the Trojans \\citep{milani93,beauge01}, and most recently, proposed to exist in the Transneptunian region \\citep{brown07}. Studies of asteroid families are very important for planetary science. Families can be used, e.g.\\ to understand the collisional history of the asteroid Main Belt \\citep{bottke05}, the outcomes of disruption events over a size range inaccessible to laboratory experiments \\citep[e.g.][]{michel03,durda07}, to understand the mineralogical structure of their parent bodies \\citep[e.g.][]{cellino02} and the effects of related dust ``showers'' on the Earth \\citep{farley06}. Obtaining the relevant information is, however, not easy. One of the main complications arises from the fact that the age of a family is, in general, unknown. Thus, accurate dating of asteroid families is an important issue in the asteroid science.  A number of age-determination methods have been proposed so far. Probably the most accurate procedure, particularly suited for young families (i.e.\\ age $<10~$Myr), is to integrate the orbits of the family members backwards in time, until the orbital orientation angles cluster around some value. As such a conjunction of the orbital elements can occur only immediately after the disruption of the parent body, the time of conjunction indicates the formation time. The method was successfully applied by \\citet{nes2002,nes2003} to estimate the ages of the Karin cluster (5.8$\\pm$0.2 Myr) and of the Veritas family (8.3$\\pm$0.5 Myr). This method is however limited to groups of objects residing on regular orbits.  For older families (i.e.\\ age $>100~$Myr), one can make use of the fact that asteroids slowly spread in semi-major axis due to the action of Yarkovsky thermal forces \\citep{farvok99}. As small bodies drift faster than large bodies, the distribution of family members in the $(a_p,H)$ plane -- where $a_p$ is the proper semi-major axis and $H$ is the absolute magnitude -- can be used as a clock. That method was used by \\citet{nes2005} to estimate ages of many asteroid families. In these estimations the initial sizes of the families were neglected, so that this methodology can overestimate the real age by a factor of as much as 1.5 -2. An improved version of this method, which accounts for the initial ejection velocity field and the action of YORP thermal torques, has been successfully applied to several families by \\citet{vok06a} and \\citet{vok06b}. Again, it is not straightforward to apply this method to families located in the chaotic regions of the asteroid belt.  \\citet{MilFar94} suggested that asteroid families, which reside in chaotic zones, can be approximately dated by \\emph{chaotic chronology}. This method is based on the fact that the age of the family cannot be greater than the time needed for its most chaotic members to escape from the family region. In its original form, this method provides only an upper bound for the age. Recently, \\citet{menios07} introduced an improved version of this method, based on a statistical description of transport in the phase space. \\citet{menios07} successfully applied it to the family of (490) Veritas, finding an age of 8.7$\\pm$1.7~Myr, which is statistically the same as that of 8.3$\\pm$0.5 Myr, obtained by \\citet{nes2003}\\footnote{Of course, Tsiganis et al.\\ (2007) and Nevorn\\'y et al.\\ (2003) used a chaotic and a regular subsets of the family, respectively.}. Despite these improvements, the chaotic chronology still suffered from two important limitations. It did not account for the variations in diffusion in different parts of a chaotic zone, which can significantly alter the distribution of family members (i.e.\\ the shape of the family). Moreover, it did not account for Yarkovsky/YORP effects, thus being inadequate for the study of older families.  In this paper we extend the chaotic chronology method, by constructing a more advanced transport model, which alleviates the above limitations. We first use the Veritas family as benchmark, since its age can be considered well-defined. Local diffusion coefficients are numerically computed, throughout the region of proper elements occupied by the family. These local coefficients characterize the efficiency of chaotic transport at different locations within the considered zone. A Monte-Carlo-type model is then constructed, in analogy to the one used by \\citet{menios07}. The novelty of the present model is that it assumes variable transport coefficients, as well as a drift in semi-major axis due to Yarkovsky/YORP effects, although the latter is ignored when studying the Veritas family. Applying our model to Veritas, we find that both (a) the shape of its chaotic component and (b) its age are correctly recovered. We then apply our model to the family of (3556) Lixiaohua, another outer-belt family but much older than Veritas and hence much more affected by the Yarkovsky/YORP thermal effects. We find the age of the Lixiaohua family to be $\\sim 155$~Myr.  We note that, depending on the variability of diffusion coefficients in the considered region of proper elements, this new transport model can be computationally much more expensive than the one applied in \\citet{menios07}. This is because, if the values of the diffusion coefficients vary a lot across the considered region, one would have to calculate them in many different points. However, even so, this computation needs to be performed only once. Then, the random-walk model can be used to perform multiple runs at very low cost, e.g.\\ to test different hypotheses about the original ejection velocities field or about the physical properties of the asteroids. On the other hand, for ``smooth\" diffusion regions in which the coefficients only change by a factor of 2-3 across the considered domain, the model can be simplified. In such regions, the age of a family can be accurately determined even by assuming an average (i.e.\\ constant over the entire region) diffusion coefficient, as we show in Section 3. ",
        "conclusions": "\\label{}  We have presented here a refined statistical model for asteroid transport, which accounts for the local structure of the phase-space, by using variable diffusion coefficients. Also, the model takes into account the long-term drift in semi-major axis of asteroids, induced by the Yarkovsky/YORP effects. This model can be applied to simulate the evolution of asteroid families, also giving rise to an advanced version of the \"chaotic chronology\" method for the determination of the age of asteroid families.  We applied our model to the Veritas family, whose age is well constrained from previous works. This allowed us to assess the quality and to calibrate our model. We first analyzed the local diffusion characteristics in the region of Veritas. Our results showed that local diffusion coefficients vary by about a factor of $2-3$ across the $(e_p,\\sin I_p)$ region covered by the (5,-2,-2) MMR.  Thus, although local coefficients are needed to accurately model (by the MCMC method) the evolution of the distribution of group-A members, average coefficients are enough for a reasonably accurate estimation of the family's age. We note though that the variable coefficients model reduces the error in $\\tau$ by $\\sim 30\\%$, but requires a computationally expensive procedure. Using the variable coefficients MCMC model, we found the age of the Veritas family to be $\\tau$ = (8.7$\\pm$1.2)~Myr; a result in very good agreement with that of \\citet{nes2003} and \\citet{menios07}.  We used our model to estimate also the age of the Lixiaohua family. This family is similar to the Veritas family in many respects; it is a typical outer-belt family of C-type asteroids, crossed by several MMRs. Like the Veritas family, only the main chaotic zone (MCZ) shows appreciable diffusion in both eccentricity and inclination. On the other hand, this is a much older family and the Yarkovsky effect can no longer be ignored. This is evident from the distribution of family members, adjacent to the MCZ. Our model suggests that the age of this family is between 100 and 230~Myr, the best estimate being $155\\pm 36~$Myr. Note that the relative error is $\\approx 23\\%$, i.e.\\ close to the $\\approx 20\\%$ that \\citet{menios07} found for the Veritas case, using a constant coefficients MCMC model.  Our model shares some similarities with the Yarkovsky/YORP chronology. Both methods are basically statistical and make use of the quasi-linear time evolution of certain statistical quantities (either the spread in $\\Delta a$ or the dispersion in $e_p$ and $\\sin I_p$), describing a family. There are, however, important differences. The Yarkovsky/YORP chronology method works better for older families and the age estimates are more accurate for this class of asteroid families (provided there are no other important effects on that time scale). On the other hand, for our method to be efficient, we need that diffusion is fast enough to cause measurable effects, but slow enough so that most of the family members are still forming a robust family structure (i.e. there is no dynamical ``sink'' that would lead to a severe depletion of the chaotic zone). Thus, our model can be applied to a limited number of families that reside in complex phase-space regions, but, in the same time, this is the only model that takes into account the chaotic dispersion of these families. There are at least a few families for which both chronology methods can be applied, thus leading to more reliable age estimates, as well as to a direct comparison of the two different chronologies. For example, the families of (20) Massalia and (778) Theobalda would be good test cases. We, however, reserve this for future work.  An important advantage of the model is that it can be used to estimate the physical properties of a dynamically complex asteroid family, provided that its age is known by independent means (e.g.\\ by applying the method of \\citet{nes2003} to the regular members of the family). A large number of MCMC runs can be performed at low computational cost, thus allowing a thorough analysis of the physical parameters of family members or the properties of the original ejection velocities field that better reproduce the currently observed shape of the family."
    },
    "0910/0910.0196_arXiv.txt": {
        "abstract": "The disruptive effect of galactic tides is a textbook example of gravitational dynamics. However, depending on the shape of the potential, tides can also become \\emph{fully} compressive. When that is the case, they might trigger or strengthen the formation of galactic substructures (star clusters, tidal dwarf galaxies), instead of destroying them. We perform $N$-body simulations of interacting galaxies to quantify this effect. We demonstrate that tidal compression occurs repeatedly during a galaxy merger, independently of the specific choice of parameterization. With a model tailored to the Antennae galaxies, we show that the distribution of compressive tides matches the locations and timescales of observed substructures. After extending our study to a broad range of parameters, we conclude that neither the importance of the compressive tides ($\\approx 15\\%$ of the stellar mass) nor their duration ($\\sim 10^7 \\U{yr}$) are strongly affected by changes in the progenitors' configurations and orbits. Moreover, we show that individual clumps of matter can enter compressive regions several times in the course of a simulation. We speculate that this may spawn multiple star formation episodes in some star clusters, through e.g., enhanced gas retention. ",
        "introduction": "That the violent collision at the origin of a galaxy merger acts to develop strong tidal features is a long-acknowledged fact of its evolution in time, as encapsulated in the Toomre sequence \\citep[see several illustrations in][]{Sandage1961, Toomre1972, Toomre1977, Hibbard1996, Laine2003, Elmegreen2007}. The burst of star formation triggered during such an event is a catalyst for the evolution of entire galactic stellar populations \\citep{Barton2000, Bridge2007, Li2008}. This process would be even more significant at high redshift ($z \\gtrsim 1$, \\citealt{Madau1996}), through either repeated accretion during the formation of a major galaxy, or enhanced merger rates driven by the formation of cosmological large-scale structures. However, the lack of resolution at these redshifts is an intrinsic difficulty when trying to identify stellar associations down to a few thousand solar masses, of a size spanning a few parsecs. At today's resolution $\\sim 0.01\\arcsec$, a full coverage of all modes of star formation triggered during a merger limits surveys to the low redshift, local Universe.  Over the past decade, high-resolution HST observations of the proto-typical merger NGC 4038/39 (the Antennae) have revealed a surprisingly high number of dense, massive young clusters of stars in its central region \\citep[][]{Whitmore1995, Meurer1995, Mengel2005, Whitmore2007}. A key question which emerges from such studies pertains to the likelihood that stellar associations will survive in the rapidly-varying background potential of the merger, once they have outlasted the early phase of internal evolution \\citep{Elmegreen1997, Bastian2006, Gilbert2007}. Based on samples of some $10^3$ Antennae clusters, \\citet{Fall2005} derived a feature-less power-law age distribution ($dN/d\\tau \\propto \\tau^{-1}$) with a median value of $\\sim 10^7 \\U{yr}$, a trend they argued to be independent of mass. Their analysis may be interpreted as implying that young clusters dissolve rapidly into their gas cloud nursery, prior to reaching virial equilibrium \\citep{Whitmore2007,deGrijs2007}. \\citet{Gieles2009} has shown that mass-dependent models of cluster evolution leading to dissolution would inevitably introduce a feature in the age distribution of Antennae clusters, lending support to the conclusions of \\citet{Fall2005}. Since cluster dissolution in equilibrium galaxies invariably lead to Gaussian distributions \\citep{Vesperini1997, Baumgardt2003}, one is drawn to ask whether the power-law properties of young cluster distributions in rapidly evolving, out-of-equilibrium mergers are shaped mainly by gravitational effects, or hydrodynamics, or both. Clearly the evolution of the large-scale gravitational potential imprints the smaller scales of star formation controlled by the hydrodynamics of accretion seeds \\citep{Klessen1998, MacLow2004}. Knowing how to connect the two would be a major step toward a global understanding of the evolution of stellar populations in mergers.  One way of addressing this problem consists in recreating observations numerically. Simulations of galaxy mergers present much more complicated kinematics that increase the complexity level of star formation modeling when compared, say, to galaxies taken in isolation. To resolve the many star-forming regions in a galaxy merger, several studies rely on smoothed particle hydrodynamics \\citep[SPH,][]{Gingold1977} simulations, assuming a star formation recipe mainly based on a local density or pressure threshold \\citep{Mihos1996, Bekki2002a, Bekki2002b, Springel2003, Li2004, Robertson2008, Karl2008, Johansson2009}. Other approaches like sticky particles \\citep[e.g.][]{Bournaud2008} or adaptive mesh refinement \\citep{Kim2009} are also employed to investigate the formation of substructures (star clusters or tidal dwarf galaxies, TDGs) at high resolution. \\citet{Barnes2004} compared the common SPH star formation rule with a new prescription based on the energy dissipation through shocks, with an application to NGC 4676 (the Mice). He showed that to gauge correctly the radiative energy from shocks and to recover a more realistic star-formation rate, it is necessary to take into account the divergence of the background velocity field, a strong hint that large-scale motion bears down on the small-scale physics of star formation. \\citet{Dobbs2009} recently showed that the surface density of neither the total gas (\\ion{H}{1} + H$_2$) nor the warm gas only (\\ion{H}{1}) were good tracers for the local star formation rate, considering an isolated disk. They also underlined that a strong spiral shock would increase the local density, but the velocity dispersion too, which preempts a one-to-one relation between star formation (averaged over the whole disk) and shocks. These results suggest to consider (on top of a density contrast) dynamical properties such as the convergence of the velocity field, either through hydrodynamical shock dissipation or through a background gravitational potential. On a smaller scale, several studies starting with \\citet{Ebert1955} and \\citet{Bonnor1956}, investigated the hydrodynamics of gaseous systems like giant molecular clouds confined by an external pressure \\citep[see e.g.][]{Kumai1986, Elmegreen1989, Harris1994, Elmegreen1997}. They showed that such an external effect could trigger the collapse of the cloud, so linking galactic-scales hydrodynamics and cluster-size phenomena. Despite this impressive body of work, the high-resolution needed to link gravitational galactic dynamics to the very small scales of star formation is still out-of-reach of current 3D numerical models.  A possible way out of this dilemma would be to impose the right boundary conditions on a small volume within which high-resolution hydro codes may resolve the physics of star formation. This paper takes a step in that direction and investigates how these boundary conditions might evolve as a function of time. \\citet{Renaud2008} already explored the evolution of the tidal field during a merger event, with a model tailored to the Antennae galaxies. They showed that \\emph{fully} compressive tides (see Section~\\ref{sec:tidaltensor}) develop during the interaction of two disk galaxies and last for few $\\times 10^7 \\U{yr}$. They concluded that such a tidal mode could play an important role in triggering the formation of star clusters, or at least prevent them from destruction. Here, we detail further and extend these results through a larger sample of progenitor galaxies and merger orbits. It is hoped that this more general yet precise description of the tidal field will in turn allow more realistic boundary conditions for volumes where star formation takes place. In Section 2, we present analytically how compressive tides may develop when two potentials overlap, as in the case of a major galaxy merger. Section 3 describes the application to $N$-body simulations and the limitations of the method. Section 4 details the Antennae reference model and Section 5 makes direct comparisons between the parameter survey and this reference. We close with a discussion on the link between the physical process described and the observational facts. ",
        "conclusions": "\\subsection{Statistics of the tidal field} The gravitationnal tidal field of galaxies is often presumed to disrupt smaller bodies orbiting in their vicinity. This point of view stems from a wealth of observational data of on-going disruption of tidal dwarf galaxies, as well as strong theoretical underpinning going all the way back to Roche and the notion of an outward stretch operating on finite-size bodies. But a close inspection of the full tidal field tensor leads one to conclude that in several realistic situations, such as in the core of flat-top density profiles or the mid-point of spiral disks, the tide will work in reverse and act to shelter structures from dissolution. Because the character of the tidal field is sensitive to the presence of clumps or other sub-structures, it seems natural to seek out quantitatively the evolution of the tides in the context of a major galaxy mergers, when large structures and smallish clumps will form in great numbers. A key question that we have addressed in this paper is whether compressive tidal modes could widely influence the early life of star clusters or TDGs, both in terms of their location in the merger, and their duration in time.  The reference Antennae model discussed earlier shows that fully compressive tidal modes spread throughout the merger, including in key areas such as the nuclei and the tidal tails. The statistics of the duration of these modes reveal that most of them last long enough to impact on the formation of bound substructures on a scale of $\\sim 200 \\U{pc}$ or smaller. This conclusion was confirmed in a parameter survey of galaxy mergers, which showed that the shape of the tidal mode distribution function is by far unaffected by details of the systems undergoing a merger (spin-orbit coupling, mass ratio, impact parameter, and so on).  Broadly speaking, compressive tidal modes always develop to reach the same level of intensity in all the simulations that we have performed. However, some trends have been highlighted. First of all, the mass fraction entering compressive modes scales with the available mass: a low-mass progenitor undergoing a collision yields a low mass-fraction of compressive tidal modes. This is clear from the spin- and mass ratio explorations. Secondly, an encounter with a small pericenter distance brings more material to the collision in a smaller volume. As a result, the number of mass elements experiencing compressive tidal modes is significantly higher for close passages.  \\subsection{Tidal field and star formation rate} For every interaction leading to mass exchange and possibly a merger, the fraction of bodies experiencing a compressive tidal field rises by a factor of at least 4, and up to 13 in the extreme case of a face-on collision, when compared with galaxies in isolation. These figures match well with the results of \\citet{diMatteo2008} who presented a study of the SFR over a large set of mergers, using both SPH and sticky particles techniques. We also remark that the starburst activity derived from their simulations would not last longer than $\\sim 500 \\Myr$, which matches well with our results. This provides indirect evidence of a close link between fluctuations in the global galactic gravitational field, and the small-scale volumes where stars are presumed to form \\citep[recall Fig.~2 of][for another, complementary argument]{Renaud2008}.  \\Citet{diMatteo2008} highlighted that retrograde encounters enhance the SFR of mergers, which is again in agreement with the evolution of the tidal field for this case (cf. models PR and RR of Section~\\ref{sec:spin}). However, they noted a negative correlation between the pericenter distance and the starburst, in the sense that close passages correspond to \\emph{low} SFR, for mergers \\emph{only}. It is yet unclear to what extent this trend can be attributed to the response of the gas in the early stages of the merger. Because the statistics reported here do not account in details for the sequence of events that take place, we can not be specific about the likelihood of forming stars at a precise time. A clearer picture will emerge once high-resolution hydrodynamics is included in the dynamical scenarios that we have reported here.  \\subsection{Age distribution function} In all merger simulations that we have carried out, a double exponential law describes well the compressive mode distribution in age, for both the longest uninterrupted sequence (\\emph{lus}) and the total time (\\emph{tt}). We note that the longest periods of the \\emph{lus} statistics correspond to orbits near the centre of the galaxies (in isolation as well as during the interaction) while the transient features that give rise to the shorter periods of the \\emph{lus} distribution were found in the merger-induced tidal structures (such as bridges and tails).  In Section~\\ref{sec:history}, we showed that the history of the gravitational tidal field along individual orbits is very complex. Indeed, to compute the net effect of the tides on a star cluster over time via a statistical approach, one should combine \\emph{both} the \\emph{lus} and \\emph{tt} distributions. This is so because the tidal field changes character (compressive or extensive) repeatedly along most orbits (see Fig.~\\ref{fig:orbits_particles}).  To illustrate this further, consider once more the case of the Antennae model. Roughly speaking, the statistics of compressive modes for that case can be split in two phases: a first phase in the time interval $\\Delta T_1 = ]0,50] \\Myr$, during which the compressive modes are mostly uninterrupted; and a second phase $\\Delta T_2 = [50,100] \\Myr$ during which the compressive modes are a collection of discrete episodes. Putting this together, a proxy for the distribution of time spent in compressive tidal modes might consist in taking the \\emph{lus} distribution over the time interval $\\Delta T_1$, followed by the \\emph{tt} statistics over the interval $\\Delta T_2$. Fig.~\\ref{fig:age} plots the \\emph{lus} and \\emph{tt} distributions (red and blue lines) showing the usual double exponential, together with the linear combination of the two described above (black curve).  \\begin{figure} \\plotone{f15.eps} \\caption{The \\emph{lus} and \\emph{tt} distributions (red and blue lines, respectively) of compressive modes for the Antennae model. The distributions consist in double exponential laws fits to the short and long timescales ($\\Delta T_1$ and $\\Delta T_2$). The black line represents a combination of both cases, using the short timescale from the \\emph{lus} and the longer one from \\emph{tt} (see text for details). The green points with horizontal binning are HST data lifted from \\citet{Fall2005}.} \\label{fig:age} \\end{figure}  If we take the bold step of identifying the duration of the modes (x-axis on Fig.~\\ref{fig:age}), with the age of stellar associations, then the vertical y-axis may be interpreted as the number of associations per age interval, $dN/d\\tau$. The black curve on Fig.~\\ref{fig:age} can be fitted with a $\\tau^{-1.08}$ power law ($\\sigma = 2\\times 10^{-4}$) over the interval $t \\in [10, 100] \\Myr$). It is striking that the duration of the compressive modes is very similar to the $\\tau^{-1}$ distribution derived from HST data by \\citet{Fall2005} for the age of young star clusters. While this does not constitute a full cause-to-effect chain of reasoning, the unexpected agreement does offer a strong hint that time-dependent tidal fields bear on the demographics of young clusters in galaxy mergers.  \\subsection{Energy of compressive modes} To characterize statistically the strength of the tides, we plot in Fig~\\ref{fig:distrib} the distribution of the maximum eigenvalue of the tidal tensor for all particles, at selected times. As mentioned before, a negative value stands for a fully compressive mode. It is clear that, at any stage of the merger, the distribution can not be fitted by simple analytical functions (such as a Gaussian or power laws). Its high degree of asymmetry with respect to the origin mainly stems from the distribution of eigenvalues in the galaxies in isolation (black curve on Fig.~\\ref{fig:distrib}), where a large fraction of the mass ($\\sim 97\\%$) experiences extensive but weak tidal modes, while the few percent mass elements nearest to the galactic center experience compressive modes of large, negative eigenvalues. The distribution of compressive and extensive modes tends to become more even-handed in the course of the merge, in terms of the occupation percentage at given $|\\lambda|$. This is illustrated at the first pericenter passage, $t = 0 \\Myr$, and at $t = 300 \\Myr$ when the merger is in full swing. One can deduce from this that the strong destructive effect of the extensive tidal field acting at some point in the system is balanced out, on average, by a compressive mode of comparable strength in some other region, like e.g. the tidal tails. Given that individual orbits enter and exit compressive modes several times during the merger (cf. Fig.~\\ref{fig:orbits_particles}) one may anticipate that transitions from compressive to extensive modes are relatively sharp over a timescale of $\\sim 2 \\Myr$.  \\begin{figure} \\plotone{f16.eps} \\caption{Normalized distribution (in percent) of the maximum eigenvalue ($\\lambda_\\mathrm{max}$) of the tidal tensor for the Antennae run. Three key steps are shown: the configuration in isolation ($t = -100 \\Myr$), the first pericenter passage ($t = 0 \\Myr$) and the merger stage ($t = 300 \\Myr$).} \\label{fig:distrib} \\end{figure}  We already noted how the full impact of the compressive tides could only be quantified once the hydrodynamics of star formation is included in the modeling. Despite this caveat, one can estimate the relative importance of compressive tides by comparing the extra binding energy they add to a cluster, to the feedback from stellar winds. For simplicity, let us assume an isotropic tidal field, i.e. described by a tensor set equal to the identity matrix times $\\lambda$. For a mass distribution of total mass $M_\\mathrm{t}$, the corresponding binding energy driven by the tide is \\begin{equation} \\label{eqn:energy} E_\\mathrm{T, cluster} = -\\frac{1}{2} \\lambda \\int_0^{M_\\mathrm{t}}{r^2 \\ dm} = -\\frac{1}{2} \\lambda \\alpha \\ M_\\mathrm{t} R_\\mathrm{t}^2 \\end{equation} where $\\alpha$ is a dimensionless quantity of the order of unity, which encapsulates details of the mass distribution, and $R_\\mathrm{t}$ is the truncation radius. Applying Eq.~\\ref{eqn:energy} to the resolution of our simulations ($M_\\mathrm{t} = 10^5 \\msun$ and $R_\\mathrm{t} = 10^2 \\U{pc}$), we get a typical value for $\\lambda$ of $\\sim -10^{-30} \\U{s^{-2}}$ ($< 0 $ for a compressive tide). This yields an estimate for the tidal binding energy of $\\sim 10^{49} \\U{ergs}$. Clearly, this value does not compare to the energy release $\\sim 10^{51} \\U{ergs}$ of a single supernova explosion, and hence SN-driven gas expulsion would proceed largely unaffected by the tides. However, the binding energy computed from Eq.~\\ref{eqn:energy} would compare favorably to the kinetic energy of an O-star's wind \\citep{Cappa1998, Martin2008}. In that case the mass loss driven by the winds of early-type stars would be slowed down considerably. The enhanced mass retention would then help stop or reduce the dissolution of young clusters, an aspect that has immediate consequences to the question related to their mass functions in mergers.  In addition to star clusters, compressive tidal modes might also be important for the formation of TDGs. Detailed high-resolution $N$-body simulations showed that the formation of TDGs is virtually impossible in purely stellar systems \\citep{Wetzstein2007}. Consequently, a dissipative component (gas) is a necessary ingredient to form TDGs. Because the compressive region will span kiloparsec-size volumes, enough stars may form that a TDG will ensue. The binding energy of the tidal field estimated from Eq.~\\ref{eqn:energy} for a mass $M_\\mathrm{t} = 10^9 \\msun$ and a size of $10 \\kpc$ (typical values for the compressive regions found in the outer parts of our merger simulations) is $\\sim 10^{57} \\U{ergs}$, or the kinetic energy release of $10^6$ type-II supernovae. Now, if one converts $10^9 \\msun$ of gas using a standard stellar initial mass function and a star formation efficiency (SFE) of $10-20 \\%$, one ends up with on order of a million type-II supernova candidate stars. While this calculation is rough, it does allow us to conclude that the energy reservoir of compressive tides must be factored in more detailed analysis of TDG formation models, as it is at least comparable to SN energy feedback.  \\subsection{Perspectives} In this last, more speculative section we explore briefly three key areas for future work that will draw from the current study:  \\begin{itemize} \\item The first and more obvious is the modeling of the time-evolution of cluster mass functions driven by galactic tides. A number of previous studies derived cluster dissolution rates based on steady external tides \\citep{Baumgardt2003, Gieles2008a, Gieles2008b, Kuepper2008, Vesperini2009, Ernst2009}. Other approaches encapsulate the tide's time-evolution in an approximate ``shock'' treatment as the model cluster orbits the host galaxy (\\citealt{Spitzer1987}; see \\citealt{Gieles2007} for a recent application to spiral galaxies). We note that the tidal tensor approach taken in this paper encompasses those mentioned here as limiting cases. But the important point to stress in our view is that the statistics of tidal fields change very significantly once we consider out-of-equilibrium dynamics. It will be important in future to take into account this evolution of the tidal field itself when considering the evolution of cluster mass distributions, especially in mergers.  \\item Secondly, the tidal field bears down on star-forming regions  at the mass- and length scales of star clusters in merging galaxies, and hence on the hydrodynamics of fragmentation and star formation in such systems. Future modeling of the first $\\sim 10^6 \\U{yr}$ of a giant molecular cloud embedded in a realistic, fully evolving background tidal field should shed new light on questions pertaining to the survival of young (open or globular) clusters \\citep{Bastian2006, Goodwin2008}. The survival of clusters is set from a combination of several effects, e.g. the expulsion of the residual gas \\citep{Goodwin1997, Geyer2001, Chen2008}; and two-body relaxation \\citep{Portegies2008}. How the rapid evolution of the tides that we have reported here changes the picture derived from the secular evolution of clusters (the weak field limit, see e.g. \\citealt{Goodwin1997, Boily2003a, Boily2003b, Baumgardt2007, Parmentier2009}) is an aspect worthy of further consideration.  \\item Finally, several authors have confirmed  in recent years the presence of multiple stellar main sequences (MS) in resolved massive clusters \\citep[see e.g][]{Piotto2007, Milone2009}. \\citet{dAntona2004} suggested that a stellar initial mass function (IMF) with a flat slope at the high-mass end would enhance CNO abundances yet leave relatively unchanged the abundances of iron-peak elements, providing a fit for the multiple threads in cluster HR diagrams. It is unclear why a bias in the stellar IMF should develop in some clusters. One possibility comes in the form of an age-spread as broad as $\\sim 10 \\Myr$, which could be interpreted as a second episode of star formation. This would open up the possibility that (i) clusters are self-enriched (by retaining slow stellar winds); or (ii) accrete gas from gas-rich environments met on their galactic orbit. In that context it is clear that compressive gravitational tides would help increase the likelihood of a second episode of star formation. Gas accretion by the cluster should impact on \\emph{both} the CNO abundances and those of the iron-peak elements, and so offers little hope of a viable explanation to this problem. Furthermore, the cluster has little time over $10 \\Myr$ to move significantly away from the volume where it first condensed before the second burst of star formation takes place. Thus, this scenario would preserve the homogeneity of the metal abundances the cluster already holds. Self-enrichment, on the other hand, should be more effective in broadening CNO abundances since the new abundances would remain correlated with those of the host cluster at birth. Recent modeling of rotating massive stars has shown that slow disk-like ejecta may be more effectively retained by the cluster than was assumed up to now \\citep{dErcole2008, Decressin2007}. Clearly the retention of light elements would be enhanced if the cluster was undergoing a compressive tidal mode at the time when such winds were operative. It is too early to pin down quantitatively the actual impact of tidal modes on such cluster evolutionary processes, which will surely require exhaustive modeling of the hydrodynamics and the chemical evolution, not done here. \\end{itemize}"
    },
    "0910/0910.0475_arXiv.txt": {
        "abstract": "{% We summarize our optical monitoring program of VY Scl stars with the SMARTS telescopes, and triggered X-ray as well as optical observations after/during state transitions of V504 Cen and VY Scl. } ",
        "introduction": "VY Scl stars are a subclass of cataclysmic variables (CVs) that are optically bright most of the time, but occasionally drop in brightness by several magnitudes  at irregular intervals (Bond 1980, Warner 1995, Honeycutt \\& Kafka 2004). The transitions between the brightness levels take place \\linebreak within days to weeks. In their high states, these variables have the largest time-averaged mass transfer rate $\\dot{M}$ (of the order of 10$^{-8}$ M$_{\\odot}$/yr) among CVs, and thus are thought to be steady accretors with hot disks. The cause of the transitions is widely debated. The conventional view is that a strong reduction (or even cessation) of the mass transfer rate, either due to the magnetic spot covering temporarily the $L_1$ region (Livio \\& Pringle 1994) or due to non-equili\\-brium effects in the irradiated atmosphere of the donor (Wu et al. 1995) causes the deep low states. A major problem of this scenario is the lack of dwarf nova (DN) outbursts during the low states (Leach et al. 1999): the disk remains subject to the thermal/viscous instability as it must drain its remaining gas onto the white dwarf (WD) through a series of DN eruptions. Another serious difficulty is that the observed dual-slope rises (Honeycutt \\& Kafka 2004), which are faster when the system is fainter, are opposite to the expected behaviour of an accretion disk since rebuilding the disk from an extended low state should initially take place on the slow viscous timescale of the low-state disk, eventually switching to a faster rise as the disk goes into DN outburst on the faster thermal timescale.  Two unconventional views invoke quasi-stable burning of the accreted hydrogen on the white dwarf surface or a magnetic nature of the white dwarf. In the former case, VY Scl stars  may be transient supersoft X-ray sources during optical low-states (Greiner \\& DiStefano 1999, Grei\\-ner \\linebreak 2000, Honeycutt 2001). Discovering more nearby supersoft X-ray sources is potentially important, since some of them may be progenitors of Type Ia supernovae; in addition, the sources may play an important role as ionizers of the interstellar medium. This picture is motivated by the behaviour of the classical supersoft X-ray source RX J0513.9--6951 which has quasi-periodic optical intensity dips of $\\sim$4 weeks duration (Southwell et al. 1996, Reinsch et al. 2000) simultaneously to X-ray high states. This is interpreted as being due to a drop in accretion rate which leads to a contraction of the WD, and in turn a hotter WD surface temperature. Besides this phenomenological similarity, there exist independent observational hints for transient supersoft X-ray emission in 5 VY Scl stars: V751 Cyg (Greiner et al. 1999), V~Sge (Greiner \\& Teeseling 1998), DW UMa (Knigge et al.\\ 2000), SW Sex (Groot et al.\\ 2001), BZ Cam (Greiner et al. 2001).  This conjecture that VY Scl stars are the low-mass extension of supersoft X-ray binaries (SSB) can only be tested by combined optical and X-ray observations. If it can be pro\\-ven, it is expected to have far-reaching consequences. For example, if a relation to SSBs can be established, it would add a whole group of optically bright objects that are much easier to study than the optically-faint SSBs. Second, because the presence of WDs in SSBs has never been proven, a connection of SSBs and VY Scl stars would add support to the standard model of supersoft X-ray sources (van den Heuvel et al.\\ 1992). Third, the irradiation of the donors in supersoft X-ray sources is much stronger than in CVs (and VY Scl stars), and therefore the mechanism proposed by Wu et al.\\ (1995) for VY Scl stars could be readily applicable to SSBs.  In the case of a magnetic primary, a low magnetic field of order 5$\\times$10$^{30}$ G cm$^3$ would disrupt the inner, otherwise unstable accretion disk, and thus prevent outbursts in low states, suggesting that VY Scl stars may all be intermediate polars (Hameury \\& Lasota 2002). While variable circular polarization (and consequently a magnetic field) has been found in two VY Scl stars (LS Peg and V795 Her), the majority has no measurable magnetic field. This conjecture can be \"easily\" tested by X-ray observations (searching for X-ray pulsations) or optical spectropolarimetry.  Here, 10 yrs after the first suggestion (Greiner \\& DiStefano 1999), we report on our results of long-term optical monitoring of southern VY Scl stars, and triggered follow-up optical and X-ray observations when a VY Scl star went into an optical low-state. ",
        "conclusions": "For V504 Cen, we tested for a hot source in March 2006. The optical spectra do not show any significant sign of HeII emission, and Chandra and XMM ToO observations do not show a strong supersoft X-ray component. The magnetic option was tested with spectro-polarimetry, but no sign of a substantial magnetic field was found. It thus appears that the suspected H surface burning during the low-states in VY Scl stars either happens at very low temperatures (below 20 eV), or is not a generic feature in this class of objects. Similarly, a magnetic nature of the WD could not be proven. This indicates that in V504 Cen both scenarios proposed as ``unconventional'' views could not be substantiated, but also the conventional Lagrange point blockage model is challenged by the very long duration of the optical low-state."
    },
    "0910/0910.5299_arXiv.txt": {
        "abstract": "{ The discovery of short-period Neptune-mass objects, now including the remarkable system HD69830 \\cite{lovis06} with three Neptune analogues, raises difficult questions about current formation models which may require a global treatment of the protoplanetary disc. Several formation scenarios have been proposed, where most combine the canonical oligarchic picture of core accretion with type I migration (e.g.~\\citealt*{terq}) and planetary atmosphere physics (e.g.~\\citealt{alib}).  To date, due in part to the computational challenges involved, published studies have considered only a very small number of progenitors at late times.  This leaves unaddressed important questions about the global viability of the models.  We seek to determine whether the most natural model -- namely, taking the canonical oligarchic picture of core accretion and introducing type I migration -- can succeed in forming objects of 10 Earth masses and more in the innermost parts of the disc. }   { This problem is investigated using both traditional semianalytic methods for modelling oligarchic growth as well as a new parallel multi-zone N-body code designed specifically for treating planetary formation problems with large dynamic range \\citep{mcn09}.} { We find that it is extremely difficult for oligarchic tidal migration models to reproduce the observed distribution.  Even under many variations of the typical parameters, including cases in which after the amount of mass in our disc is greatly increased above the standard Hayashi minimum-mass model, we form no objects of mass greater than 8 Earth masses.  By comparison, it is relatively straightforward to form icy super-Earths.}   {We conclude that either the initial conditions of the protoplanetary discs in short-period Neptune systems were substantially different from the standard disc models we used, or there is important physics yet to be understood and included in models of the type we have presented here.} ",
        "introduction": "\\label{section:intro}  At present, there are $\\sim\\!$19 known extrasolar planets\\footnote{Data taken from the Extrasolar Planets Encyclopedia, Schneider, J., http://exoplanet.eu} with estimated masses between 0.03 and 0.12 \\MJup, or between $\\sim\\!10$ and $\\sim\\!40$ \\ME.  With one exception -- OGLE-05-169L b, at 2.8 AU -- all of the planets have semimajor axes smaller than 1 AU, and so are reasonably called hot Neptunes.  In fact, with only one more exception, HD69830 d, all have semimajor axes $< 0.23$ AU.  To determine whether the objects are genuine Neptunes, i.e.~ice giants, and not merely very large rocky bodies better thought of as super-Earths, knowledge of the mean densities is required.  Three of these objects have known radii, and two have mean densities compatible with being a Neptune-like body. GJ436 b \\citep{butler2004} has a density of $\\sim\\!1.69\\,\\mathrm{g/cm^3}$ \\citep{torres}, and HAT-P-11 b \\citep{bakos} has a density of $\\sim\\!1.33\\,\\mathrm{g/cm^3}$; compare Neptune with $1.64\\,\\mathrm{g/cm^3}$.  The HD69830 system \\citep{lovis06} is particularly interesting.  It contains three Neptune-like objects: HD69830 b at 0.0785 AU, of 10.5 \\ME$\\,$ and eccentricity 0.1; HD69380 c at 0.186 AU, of 12.1 \\ME$\\,$ and eccentricity 0.13; and HD69380 d at 0.63 AU, of 18.4 \\ME, and eccentricity 0.07 ($\\pm0.07$, so a near-circular orbit is possible). The system was also reported to have a disc of warm infrared-emitting dust located between planets c and d \\citep{beichman2005}, presumably due to a collisionally active remnant asteroid belt.  More recent work \\citep{lisse07} suggests instead that the infrared excess is due to a debris disc outside the known planets at $\\simeq 1$ AU, probably resulting from the breakup of an asteroid.  This system clearly provides a rigorous test of planet formation and migration theories.  Indeed, short-period Neptune-mass bodies have several advantages as probes of planetary origins.  They are large enough to be observable, but small enough that we need not consider gravitational disc instability as a formation process.  They are also large enough to undergo significant type I migration, but not so large that they can open a gap in the disc and undergo type II migration \\citep{pap84, bryden99, crida06}.  Accordingly, hot Neptunes can yield less ambiguous tests of type I migration than more massive planets which could have significantly perturbed the gas disc they were embedded in, and their existence provides strong evidence that some kind of type I migration is operating in protoplanetary discs.  Hot Neptunes are therefore useful for exploring the intersection of the classical oligarchic core accretion picture (\\citealt*{ko98}; e.g.~\\citealt*{cham01, thommes03}) with type I migration \\citep{ward97}.  Two main classes of scenario exist in the literature: one which concentrates on the atmospheric gas physics, and one which concentrates on the dynamics of the protoplanetary interactions.  \\cite{alib} incorporate sophisticated atmospheric physics and follow the evolution of effectively isolated cores through the disc as they grow via planetesimal accretion, migrate inwards due to type I effects, accrete gas after the accretion rate of solids drops, and finally have their atmospheric mass reduced after arriving in the short-period region by evaporation.  However, the history for the HD69830 system proposed in \\cite{alib} is difficult to reconcile with the oligarchic paradigm \\citep{ko98}.  Their best-fitting model involves exactly three seed objects of masses $M = 0.6\\,\\ME$ at semimajor axes 3 AU, 6.5 AU, and 8 AU, which migrate to small semimajor axis through a mostly pristine planetesimal disc and therefore can accrete a fair amount of material \\citep{growthofmig}. Where did these three progenitor objects come from?  In an oligarchic framework, accretion in a narrow region -- even in the presence of type I drag (see \\citealt{komI, mcn05, komII}) -- results in many roughly equal-mass bodies separated by a distance of $10\\sim20$ Hill radii.  These bodies then interact and merge in the giant-impact phase of planet formation, and consume the accessible remnant planetesimal disc.  That is, if a half-Earth-mass seed can successfully form at 3 AU, there should be a large number of similar seeds of varying mass inside, each of which will accrete and migrate in similar ways. Moreover, the interior seeds are likely to have formed first.  The net result is that any such inward-migrating seed should be migrating through a region highly depleted by the previous seeds, and by the time an 0.6 \\ME$\\,$ core is formed at 3 AU much of the interior disc should have gone to completion, and possibly formed its own planets \\citep{cham08}. There is no obvious mechanism to suppress embryo growth everywhere in the disc except at three specific locations.  An additional problem is that simulations have demonstrated that cores migrating through a planetesimal swarm are unable to grow at the rate prescribed by the \\cite{alib} model.  In the presence of gas drag, mean motion resonances cause the majority of the interior planetesimals to be shepherded rather than being accreted, resulting in planets whose masses are too low, and in the wrong ratio, compared to those observed in the HD69830 system \\citep{payne2009}.  Though not aiming at HD69830 in particular, \\cite{terq} study the formation of hot super-Earths and Neptunes by following the evolution of $10-25$ planets of $0.1$ or $1\\,\\ME$ placed interior to 2 AU under type I drag, include tidal interactions with the star, and use an inner cavity in the gas disc at $\\sim\\!0.05$ AU.  For various disc parameters, they succeed in making several objects of mass $\\geq 8 \\,\\ME$, with a maximum of 12 \\ME.  They find that the typical result has between 2--5 planets, usually on near-commensurable orbits (with strict commensurability often being broken by tidal circularization). It is not clear whether this will scale up to larger masses, as they performed only one run with total mass larger than 12 \\ME (namely 25 \\ME), and so an HD69830-like system would not appear in their results even if the model would have succeeded in producing them.  Very little material was lost from their runs, suggesting it might be possible. However, their initial configurations are difficult to reconcile with an oligarchic migration process.  Scenarios involving oligarchic formation and type I migration do not exhaust the possible formation histories of the low-mass hot exoplanet population, although they are arguably the most natural. \\cite{raymond08} surveys several other proposed possibilities (for the terrestrial-mass regime): in situ formation; shepherding by migrating giant planets or secular resonance sweeping; tidal circularization; and photoevaporation of giant planets.  We will concentrate instead on the simplest oligarchic type I migration picture, which should have more success self-consistently generating an icy planet population, and attempt to determine whether the fiducial models can reproduce the observed distribution of hot Neptunes.  If they can, all the better; if they cannot, then the specifics of their failure may point in the direction of a solution and help us choose between the various possibilities on offer.  To understand how short-period Neptunes are formed, we need to move toward self-consistent global N-body models such as those which have proved useful in understanding the formation of the terrestrial planets and the outer solar system.  One problem presents itself immediately: the short-period exoplanet problem has a formation time to dynamical time ratio $\\sim\\!30$ times larger than the equivalent terrestrial formation problem, and $\\sim\\!1000$ times that of the equivalent Jovian core formation problem.  The lifetime of the gas disc (with a probable upper limit of $\\sim6$ Myr) is a common reference time-scale for each problem, as gas giants must form while there is still gas for them to accrete, and even ice giants which attain short-period orbits must have migrated while there is gas present.  However, the characteristic orbital periods for each formation problem vary from $0.01$ yr at 0.05 AU to $0.35$ yr at 0.5 AU to $11$ yr at 5.0 AU.  (Indeed, $0.01$ yr is optimistic, as many exoplanets are closer to their parent stars than 0.05 AU.)  Since most state-of-the-art planetary N-body codes require the integration timestep to be some small fraction of the orbital period of the innermost object, preserving the same wall-clock run times that researchers have become accustomed to would usually require reducing the number of planetesimals in a simulation by orders of magnitude, producing an unacceptable decrease in resolution.  It is extremely unlikely that this difficulty can be eliminated entirely, as it is a consequence of the strong dependence of orbital period on semimajor axis in the Kepler problem.  That said, the challenge can be managed to some extent by taking advantage of the scale separation caused by the troublesome dynamic range.  The standard approaches use a common drift timestep for all particles (and therefore ``over-integrate'' the more distant objects) and compute the (non-encountering) forces between the particles at the same frequency, both of which involve far more computation on the distant, slow-moving objects than is necessary to preserve qualitatively accurate dynamics. In \\cite{mcn09} the authors combine various techniques in the literature (\\citealt*{dll98, cham99, saha92}) to construct a new algorithm which allows for radial zones with different timesteps and different inter-zone force evaluation frequencies, but reduces to the proven techniques of \\cite{dll98} and \\cite{cham99} for objects within the same zone.  This allows new trade-offs between force accuracy and run time.  In our first paper, \\cite{mcn09}, we addressed the numerical challenges of studying oligarchic models of short-period exoplanet formation.  Here we apply a new code with a parallel implementation of those methods, to determine the ``reference population'' of planets predicted by the fiducial models.  We integrate global planetesimal discs extending from 0.05 AU to 10 AU under various models of the protoplanetary disc.  Our models, both semianalytic and N-body, are described in \\S\\ref{modeldescr}, and the simulation conditions in \\S\\ref{simcond}.  Simulation results are presented in \\S\\ref{results} and discussed in \\S\\ref{discussion}, and we conclude in \\S\\ref{section:conc}. ",
        "conclusions": "\\label{section:conc}  We performed 48 simulations of various oligarchic migration scenarios to determine whether the simplest standard approach can succeed in forming a population of short-period Neptune systems under common assumptions for the protoplanetary disc parameters. Multiple numerical techniques were applied: semianalytic techniques for the first 0.4 Myr, our new parallel multizone N-body code for the accretion phase while gas is present up to 6 Myr, and the more traditional SyMBA approach for the late stage to 100 Myr.  We find that over a wide range of disc conditions, it is difficult to form planets of mass greater than $3-4 \\ME$.  Our most successful runs involved $\\sim$ 5 times the mass of the MMSN, surface density varying as $r^{-1/2}$, a disc decay time-scale of 1 Myr, and a migration efficiency of 0.3.  Our most common planet outcomes are of Earth-mass objects, with the terrestrial planets having ice fractions from 0.0 to 0.75 (the maximum possible in our simulations).  The larger objects have higher ice fractions, with the median being 0.60 for objects above 1 \\ME, and 0.36 below.  In none of the cases did we succeed in forming an object of greater than 7.5 \\ME inside 2 AU, much less inside 0.5 AU, and the total embryo mass remaining inside 2 AU was always less than 17 \\ME.  The existence of an upper limit and the weak dependence on most parameters is in accordance with the predictions of \\cite{komII}, and in rough agreement with the predictions of \\cite{cham08} except at large disc masses.  Nevertheless, we should be wary of making predictions based on these results regarding extrasolar planetary systems, as they entirely fail to reproduce the short-period Neptune planetary population that we know exists.  Our failure can be compared to several previous successes in the literature, which either (1) adopt initial conditions which are not easily reconciled with an oligarchic growth picture, (2) use an inner edge to the migration (which is defensible but will have difficulty explaining more distant Neptunes), or (3) neglect inter-embryo dynamics and use an embryo merger condition which is calibrated to an effective inter-embryo separation (10 Hill radii) which is considerably smaller than we observe.  Varying parameters which we kept constant such as the gas-to-dust ratio, incorporating additional accretion physics such as fragmentation, moving to extremely large disc masses or extremely weak migration, and simply performing more runs (and hoping for fortuitous late mergers), could possibly succeed in improving the maximum mass reached by a factor of two and therefore into the Neptune-like region.  It seems quite unlikely that they will increase the median mass enough to comfortably produce a population of multiple-planet short-period Neptune systems.  We conclude that forming a system like HD69830 will probably require a significant revision to the simple models explored here.  If the standard oligarchy-plus-type-I-migration picture fails to reproduce the observed distribution of short-period exoplanets even at more extreme parameter values, then we must consider non-standard models.  Oligarchy is relatively well understood both analytically and numerically; by comparison type I migration is sensitive to poorly understood properties of the gas disc such as disc turbulence and local thermodynamic time-scales.  In a follow-up paper we consider the implications for hot exoplanet formation via oligarchy of alternate migration models (such as \\citealt*{paard06}) which show some promise."
    },
    "0910/0910.0643_arXiv.txt": {
        "abstract": "% The {\\sl F-GAMMA}-project is the coordinated effort of several observatories to understand the AGN phenomenon and specifically blazars via multi-frequency monitoring in collaboration with the {\\sl Fermi}-GST satellite since January 2007. The core observatories are: the Effelsberg 100-m, the IRAM 30-m and the OVRO 40-m telescope covering the range between 2.6 and 270 GHz. Effelsberg and IRAM stations do a monthly monitoring of the cm to sub-mm radio spectra of 60 selected blazars whereas the OVRO telescope is observing roughly 1200 objects at 15 GHz with a dense sampling of 2 points per week. The calibration uncertainty even at high frequencies, is of a few percent. 47\\% of the Effelsberg/Pico Veleta sample is included in the LBAS list. An update of the monitored sample is currently underway. ",
        "introduction": "The {\\sl unified model} paradigm for the Active Galactic Nuclei (AGN), attributes the wealth of AGN classes mostly to different parameters of the same principal system viewed at different angles of the line-of-sight to the relativistic jet. Radiogalaxies and blazars are two extremes: in the former the jet is seen from large whereas in the latter from small angles ($\\theta\\le 20^\\circ$).   Blazars are dominated by relativistic beaming and they exhibit intense variability at all wavelengths and time scales, highly superluminal apparent speeds, a significant degree of fractional polarization and polarization variability. Shock-in-jet  \\citep[][]{marscher1985, hughes1985, marscher1996}, colliding relativistic plasma shells \\citep{guetta2004}, lighthouse effect \\citep{camenzind1992} or MHD-instabilities in the accretion disks \\citep{begelman1980,villata1999}  are some of the models that attempt to explain variability at different time scales. Although the exact physical processes at play are unclear, the study of the temporal behavior of the SED can shed light on emission mechanism since different mechanisms predict different variability patterns \\citep[e.g.][]{bottcher2002}. Hence, multi-frequency monitoring of blazars is essential in understanding the blazar physics. ",
        "conclusions": "{\\bf The F-GAMMA alliance: }The Large Area Telescope (LAT) on-board {\\sl Fermi}-GST offers a unique opportunity covering the $4\\pi$ sky every three hours providing $\\gamma$-ray light curves at $\\sim$30 MeV - 300 GeV. To exploit that, we have initiated a tightly coordinated ``alliance'' of teams \\citep[F-GAMMA project,][]{angelakis2008, fuhrmann2007} in order to understand the AGN physics via multi-frequency monitoring. The core program involves the 100-m MPIfR telescope (Effelsberg, Bonn, 2.64 - 43.00 GHz), the 30-m IRAM telescope (Pico Veleta, Granada, 86, 142, 220 and 270 GHz) and the 40-m OVRO telescope (Owens Valley, California, 15 GHz). The Effelsberg and the IRAM telescopes do a monthly monitoring of $\\sim$60 blazars compiled on the basis of their $\\gamma$-ray detectability \\citep[EGRET][]{hartman1999} and their presence in the ``high priority list'' of the LAT AGN group. The OVRO telescope is monitoring a statistically complete sample of 1200 sources based on the CGRaBS catalog \\citep{healy2008}. The sampling rate is 2-3 times a week \\citep[see][]{richards2009,maxmoerbeck2009}. \\\\{\\bf The first 2.5 years:} The F-GAMMA project has been running since January 2007 at Effelsberg and June 2007 at OVRO and Pico Veleta. The data products of the first 2.5 years can be publicly accessed at {\\tt www.mpifr.de/div/vlbi/fgamma}. The cross-correlation of the Effelsberg/Pico sample with the LBAS list \\citep{abdo2009} showed that 47\\% of our sources are detected by {\\sl Fermi}-GST. A $\\chi^2$ test shows that practically all our sources are variable at all frequencies at a confidence level of 99.9\\,\\%. The variability amplitude (as parameterized by the modulation index $m=rms/<S>$ ) increases with frequency as expected from frequency flare evolution arguments. For Effelsberg, the calibration uncertainties are kept to the level of $3-5\\,\\%$ at high frequencies ($\\nu\\ge 23.05$\\, GHz) and of $0.5 - 1\\,\\%$ at the low frequencies ($\\nu<23.05$\\, GHz). The characteristic time scales present in the light curves vary from weeks to years while there is a wealth of variability patterns both in the lightcurves and in the spectra. The findings are summarized in Fuhrmann et al. (in prep.), Richards et al. (in prep.) and Angelakis et al. (in prep.). Given the limited fraction of our sources in the LBAS list we are currently compiling an updated sample."
    },
    "0910/0910.5120_arXiv.txt": {
        "abstract": "We describe the spectroscopic target selection for the Galaxy And Mass Assembly (GAMA) survey. The input catalogue is drawn from the Sloan Digital Sky Survey (SDSS) and UKIRT Infrared Deep Sky Survey (UKIDSS).  The initial aim is to measure redshifts for galaxies in three $4\\times12$ degree regions at 9\\,h, 12\\,h and 14.5\\,h, on the celestial equator, with magnitude selections $r<19.4$, $z<18.2$ and $K_{\\rm AB}<17.6$ over all three regions, and $r<19.8$ in the 12-h region. The target density is $1080\\,\\persqdeg$ in the 12-h region and $720\\,\\persqdeg$ in the other regions.  The average GAMA target density and area are compared with completed and ongoing galaxy redshift surveys.  The GAMA survey implements a highly complete star-galaxy separation that jointly uses an intensity-profile separator ($\\dsg=r_{\\rm psf}-r_{\\rm model}$ as per the SDSS) and a colour separator.  The colour separator is defined as $\\dsgjk=J-K-f(g-i)$, where $f(g-i)$ is a quadratic fit to the $J-K$ colour of the stellar locus over the range $0.3<g-i<2.3$.  All galaxy populations investigated are well separated with $\\dsgjk>0.2$.  From two years out of a three-year AAOmega program on the Anglo-Australian Telescope, we have obtained 79\\,599 unique galaxy redshifts.  Previously known redshifts in the GAMA region bring the total up to 98\\,497.  The median galaxy redshift is 0.2 with 99\\% at $z<0.5$. We present some of the global statistical properties of the survey, including \\newtext{$K$-band galaxy counts}, colour-redshift relations and preliminary $n(z)$. ",
        "introduction": "\\label{sec:intro}  Galaxy redshift surveys provide a fundamental resource for studies of galaxy evolution. The redshift of a galaxy can be used to obtain a distance assuming a set of cosmological parameters, modulo peculiar velocities, and a well-defined selection function enables the comoving number density of galaxies to be estimated as a function of various properties, e.g., galaxy luminosity functions \\citep{Schechter76,BST88,norberg02,blanton03ld}.  In addition, using the combined sky distribution and distance information, the clustering properties of galaxies can be determined \\citep{DGH78,dLGH88,norberg02clus,zehavi05} and the velocity dispersion of galaxies in groups and clusters can be used to infer dark-matter halo masses \\citep{Zwicky37,HG82,MFW93,carlberg96,eke04,berlind06}.  The target selection algorithm and area covered by a redshift survey relate to the redshift range and volume surveyed.  The industry of these surveys started in the 1980's with surveys of $\\sim2500$ galaxies over large sky areas \\citep{davis82,saunders90} and a deeper survey of 330 galaxies over $70\\,\\sqdeg$ \\citep{peterson86}. It expanded and diversified in the 1990's with surveys such as the wide-but-shallow CfA2 redshift survey, Las Campanas Redshift Survey, ESO Slice Project, and the deep-but-narrow Canada-France Redshift Survey.  Figure~\\ref{fig:compare-grs} shows the surface density of galaxy spectra versus area for these and other surveys, and Table~\\ref{tab:z-survey-list} gives selections and references. The target density is a wavelength-independent metric for depth, at least for high-completeness magnitude-limited surveys. The advent of multi-object spectrographs such as the Two-Degree Field (2dF; \\citealt{lewis02df}) and Sloan Digital Sky Survey (SDSS; \\citealt{york00}) telescope have enabled redshift surveys of $>10^5$ galaxies: the 2dF Galaxy Redshift Survey (2dFGRS) and SDSS Main Galaxy Sample (MGS).  \\begin{figure} \\centerline{ \\includegraphics[width=\\singlecolsize\\textwidth]{f01.ps} } \\caption{Comparison between field galaxy surveys with spectroscopic redshifts: {\\it squares} represent predominantly magnitude-limited surveys; {\\it circles} represent surveys involving colour cuts for photometric redshift selection; while {\\it triangles} represent highly targeted surveys. The colours represent different principal wavelength selections as in the legend. Filled symbols represent completed surveys. See Table~\\ref{tab:z-survey-list} for survey names and references.} \\label{fig:compare-grs} \\end{figure}  \\begin{table*} \\caption{List of field galaxy redshift surveys. The surveys shown in Fig.~\\ref{fig:compare-grs} are listed in order of increasing area. They are mostly magnitude limited galaxy samples except for some with colour selection (CS). The information was obtained from the references and the survey websites.} \\label{tab:z-survey-list} \\begin{tabular}{lllcl} \\hline abbrev.  & survey name  &  selection(s) & area/$\\sqdeg$ & reference \\\\ \\hline CFRS     & Canada-France Redshift Survey & $I_{\\rm AB}<22.5$ & $0.14$ & \\citealt{lilly95} \\\\ LBG-z3   & Lyman Break Galaxies at $z\\sim3$ Survey & $R_{AB}<25.5$ with CS$^a$ & $0.38$ & \\citealt{steidel03} \\\\ VVDS-deep& VIMOS VLT Deep Survey deep sample & $I_{\\rm AB}<24.0$ & $0.5$ & \\citealt{lefevre05} \\\\ CNOC2    & Canadian Network for Obs.\\ Cosmology 2 \\rlap{...} & $R<21.5$ & $1.5$ & \\citealt{yee00} \\\\ zCOSMOS  & Redshifts for the Cosmic Evolution Survey & $I_{\\rm AB}<22.5$, $I_{\\rm AB}\\la24$ with CS$^b$  & $1.7$ & \\citealt{lilly07} \\\\ DEEP2    & Deep Evolutionary Exploratory Probe 2 \\rlap{...} & $R_{\\rm AB}<24.1$ with CS$^c$ & $2.8$ & \\citealt{davis03} \\\\ Autofib  & Autofib Redshift Survey & $b_{J}<22.0$ & $5.5$ & \\citealt{ellis96} \\\\ H-AAO    & Hawaii+AAO K-band Redshift Survey & $K<15.0$ & $8.2$ & \\citealt{huang03} \\\\ AGES     & AGN and Galaxy Evolution Survey & incl.\\ $R<20.0$, $B_{\\rm W}<20.5$ & $9.3$ & \\citealt{watson09} \\\\ VVDS-wide& VIMOS VLT Deep Survey wide sample & $I_{\\rm AB}<22.5$ & $12.0$ & \\citealt{garilli08}\\\\ ESP      & ESO Slice Project & $b_{J}<19.4$ & $23.3$ & \\citealt{vettolani97}\\\\ MGC      & Millennium Galaxy Catalogue & $B<20.0$ & $37.5$ & \\citealt{liske03}\\\\ {\\bf GAMA}&Galaxy And Mass Assembly Survey & $r<19.8$, $z<18.2$, $K_{\\rm AB}<17.6$ & $144$ & --- this paper --- \\\\ 2SLAQ-lrg& 2SLAQ Luminous Red Galaxy Survey & $i<19.8$ with CS$^d$ & $180$ & \\citealt{cannon06}\\\\ SDSS-s82 & SDSS Stripe 82 surveys & incl.\\ $u\\la20$, $r<19.5$ with CS$^e$ & $275$ & \\citealt{sdssDR4}\\\\ LCRS     & Las Campanas Redshift Survey & $R<17.5$ & $700$ & \\citealt{shectman96}\\\\ WiggleZ  & WiggleZ Dark Energy Survey & ${\\rm NUV}<22.8$ with CS$^f$ & $1000$ & \\citealt{drinkwater09}\\\\ 2dFGRS   & 2dF Galaxy Redshift Survey & $b_{J}<19.4$ & $1500$ & \\citealt{colless01}\\\\ DURS     & Durham-UKST Redshift Survey & $b_{J}<17.0$ & $1500$ & \\citealt{ratcliffe96}\\\\ SAPM     & Stromlo-APM Redshift Survey & $b_{J}<17.1$ (1 in 20 sampling) & $4300$ & \\citealt{loveday92}\\\\ SSRS2    & Southern Sky Redshift Survey 2 & $B < 15.5$ & $5500$ & \\citealt{dacosta98}\\\\ SDSS-mgs & SDSS Main Galaxy Sample & $r<17.8$ & $8000$ & \\citealt{strauss02}\\\\ SDSS-lrg & SDSS Luminous Red Galaxy Survey & $r<19.5$ with CS$^g$ & $8000$ & \\citealt{eisenstein01} \\\\ 6dFGS    & 6dF Galaxy Survey & $K<12.7$, $b_{J},r_F,J,H$ limits & $17000$ & \\citealt{jones09} \\\\ CfA2     & Center for Astrophysics 2 Redshift Survey & $B<15.5$ & $17000$ & \\citealt{falco99}$^h$ \\\\ PSCz     & IRAS Point Source Catalog Redshift Survey & $60\\mu m_{\\rm AB}<9.5$ & 34000 & \\citealt{saunders00} \\\\ 2MRS     & 2MASS Redshift Survey & $K<12.2$ & 37000 & \\citealt{erdogdu06} \\\\ \\hline \\end{tabular} \\begin{flushleft} Notes: $^a$CS by $U$-band `dropouts' for photometric redshifts ($z_{\\rm ph}$) $\\sim2.5$--3.5; $^b$CS for $z_{\\rm ph}\\sim1.4$--3.0, deeper limit over $1\\,\\sqdeg$; $^c$CS for $z_{\\rm ph}\\ga0.7$; $^d$CS for $z_{\\rm ph}\\sim0.45$--0.8; $^e$CS for $z_{\\rm ph}\\la0.15$; $^f$CS by ${\\rm FUV-NUV}>1.5$ (GALEX bands) and $20.5<r<22.5$ for $z_{\\rm ph}\\sim0.5$--1.0; $^g$CS for $z_{\\rm ph}\\sim0.2$--0.5; $^h$Reference is for the Updated Zwicky Catalog that includes CfA2 redshifts. \\end{flushleft} \\end{table*}  The Galaxy And Mass Assembly (GAMA) project has at its core a galaxy redshift survey using the upgraded 2dF instrument AAOmega on the Anglo-Australian Telescope (AAT). GAMA will eventually incorporate a range of new surveys from UV, visible, IR and radio wavelengths \\citep{driver09}.  The redshift survey uses for its input catalogue data from the Sloan Digital Sky Survey and United Kingdom InfraRed Telescope (UKIRT).  The primary goals of the redshift survey are measurement of the halo mass function \\citep{eke06}, galaxy stellar mass function \\citep{cole01}, and the merger rates of galaxies \\citep{depropris05}: for systems of the lowest possible masses (at low redshift $z<0.05$), and for their evolution out to $z\\sim0.5$.  In terms of depth and area of magnitude-limited surveys (squares in Fig.~\\ref{fig:compare-grs}), GAMA bridges the gap between the wide but shallower surveys like 2dFGRS and SDSS MGS, and the deep but narrower surveys such as those using the VIsible MultiObject Spectrograph (VIMOS) on the Very Large Telescope (VLT).  The outline of the paper is as follows.  The imaging data, magnitude measurements and initial catalogues are described in \\S~\\ref{sec:imaging}. The target selection is described in \\S~\\ref{sec:target-select}: star-galaxy separation, magnitude limits, and other quality checks.  The pre-existing and GAMA spectroscopic data sets are outlined in \\S~\\ref{sec:spec}. An analysis of results as pertaining to the star-galaxy separation and other selection criteria is presented in \\S~\\ref{sec:results}. In other survey papers, the scientific and multi-wavelength database aims are described in \\citet{driver09}, and the tiling strategy is described in \\citet{robotham09}.  Magnitudes are corrected for Milky-Way extinction using the dust maps of \\citet{SFD98} except for fibre magnitudes. Extinction in the bands $u,g,i,z,J,K$ are obtained from SDSS $r$-band extinction using fixed ratios (1.873864, 1.378771, 0.758270, 0.537623, 0.323, 0.131).\\footnote{The SDSS extinction ratios are given in table~22 of \\citet{stoughton02}.  The ratios for $J$- and $K$-band extinctions were obtained from UKIRT WFCAM science archive \\citep{hambly08} data matched to SDSS $r$-band extinction.}  The UKIRT magnitudes are converted to the AB system using $J_{\\rm AB} = J + 0.94$ and $K_{\\rm AB} = K + 1.90$ \\citep{hewett06}.  The contours used to represent bivariate distributions (Figs.~\\ref{fig:sg-sep-histo}--\\ref{fig:test-auto-jk}, \\ref{fig:color-bias}--\\ref{fig:size-z}, \\ref{fig:obs-color-z}) are logarithmically spaced in number density, with four levels per factor of ten. ",
        "conclusions": "\\end{equation} Only objects satisfying these criteria are targeted in the main survey.  The GAMA UKIDSS selection was based on non-pipeline \\textsc{SExtractor} magnitudes. Thus, the final star-galaxy separation (Eq.~\\ref{eqn:ukidss-sdss-star-gal-sep}) was determined using \\textsc{auto} mags for $J-K$ and SDSS model mags for $g-i$ (data from Stripes 9--12). \\newtext{Figure~\\ref{fig:test-auto-jk} shows histograms in $\\dsgjk$ using these magnitudes. In order to test the position of the stellar locus, unresolved samples ($\\dsg < 0.05$; not used for targeting) were selected in separate $1\\degr \\times 1\\degr$ regions (grey lines in Fig.~\\ref{fig:test-auto-jk}). The median value of $\\dsgjk$ within each region varied from $-0.05$ to 0.04 with 90 per cent of the region values between $-0.03$ and 0.02. This demonstrates that the stellar locus fit applies to \\textsc{auto} mags equally well.}  \\begin{figure} \\centerline{ \\includegraphics[width=\\singlecolsize\\textwidth]{f08.ps} } \\caption{\\newtext{Histograms in $J-K$ star-galaxy separation parameter using \\textsc{SExtractor} \\textsc{auto} magnitudes.  The histograms represent UKIDSS-SDSS matched samples divided into unresolved (not selected for targeting), marginally resolved (selected if $\\dsgjk>0.2$) and strongly resolved (selected regardless of $\\dsgjk$) in SDSS $r$-band imaging. The thin grey line histograms represent unresolved samples in 36 randomly selected $1\\degr \\times 1\\degr$ regions.}} \\label{fig:test-auto-jk} \\end{figure}  \\subsection{Magnitude limits} \\label{sec:mag-limits}  The main scientific goal of GAMA that drives the choice of the minimum width of the survey geometry, and the magnitude selection, is the measurement of the halo mass function \\citep{driver09}.  We chose $r$-band selection because it is most directly correlated with spectral S/N obtained (the filter falls in the middle range of the spectrograph). This ensures a high redshift success rate for a given target density. The $r$-band limits were chosen to give an average target density up to an order of magnitude higher than the SDSS MGS ($90\\,\\persqdeg$) and 2dFGRS ($140\\,\\persqdeg$). Given the limitations of efficient observing over two or three lunations each year, three fields were chosen covering 6 hours in RA. We compromised between area and depth by choosing a limit of $r<19.4$ in G09 and G15 ($670\\,\\persqdeg$), and $r<19.8$ in G12 ($1070\\,\\persqdeg$). These were defined using Petrosian magnitudes, following the strategy of the SDSS MGS.  In consideration of measuring the stellar mass function, we included a near-IR selection using SDSS $z$-band and UKIDSS $K$-band.  To ensure reliability and reasonable redshift success rate, these were also constrained by an $r$-band selection ($r_{\\rm model}<20.5$).  The choice of SDSS model magnitudes rather than Petrosian is a consequence of the noise statistics. For Petrosian magnitudes, the noise is well behaved to $r\\simeq20$ \\citep{stoughton02}, while for fainter objects the model magnitudes are more reliable. Figure~\\ref{fig:mag-test} shows the pipeline-output magnitude errors versus magnitude.  Also, the $K$-band selection was based on \\textsc{auto} magnitudes, and both \\textsc{auto} and model magnitudes use elliptical apertures.  The additional selections were a small sample to $z_{\\rm model}<18.2$ and a sample to $K_{\\rm AB,auto}<17.6$.  \\begin{figure} \\centerline{ \\includegraphics[width=\\smallcolsize\\textwidth]{f09.ps} } \\caption{Magnitude errors versus magnitude. The solid lines show the median errors obtained from the SDSS catalogue, with the regions representing the inter-quartile range (top for Petrosian, lower for model magnitudes). The vertical dash-dotted lines represent the $r$-band limits used in this paper (19.4 and 19.8 using Petrosian, and 20.5 using model magnitudes).} \\label{fig:mag-test} \\end{figure}  Within the GAMA regions, the main survey selections are given by: \\begin{equation} \\begin{array}{ll} r_{\\rm petro}<19.4 & \\mbox{\\textsc{or}} \\\\ r_{\\rm petro}<19.8 \\mbox{~~\\textsc{and}~ in the G12 area} & \\mbox{\\textsc{or}} \\\\ z_{\\rm model}< 18.2 \\mbox{~~\\textsc{and}~~} r_{\\rm model} < 20.5 & \\mbox{\\textsc{or}} \\\\ K_{\\rm AB,auto}  < 17.6 \\mbox{~~\\textsc{and}~~} r_{\\rm model} < 20.5  \\: . \\end{array} \\label{eqn:mag-limits} \\end{equation} Including the near-IR selections increases the G12 target density marginally (to $1080\\,\\persqdeg$) while increasing the G09 and G15 target density to $720\\,\\persqdeg$. Figure~\\ref{fig:color-bias} shows the colour bias for the near-IR selections.  The $z$-band selection is complete to $(r-z)_{\\rm model} < 2.3$ at the faint limit, while the $K$-band selection is complete to $r_{\\rm model} - K_{\\rm AB,auto} < 2.9$ at the faint limit. A $z_{\\rm model}<18.2$ selection is formally missing 0.3\\% of objects because of the $r_{\\rm model}$ limit, while a $K_{\\rm AB,auto}<17.6$ selection is formally missing about 1\\% of objects. This is after applying star-galaxy separation. However, only very red objects are missed, which are more likely to be stars in spite of the star-galaxy separation or have incorrectly measured colours caused by mismatched apertures in the case of $r-K$ (the practical impact of these joint limits is discussed later in \\S~\\ref{sec:z-distribution}).  \\newtext{See also Fig.~\\ref{fig:K-counts}, the $r<20.5$ limit only makes an obvious impact in the galaxy number counts at $K_{\\rm AB,auto}>17.8$ that is above our selection limit.}  \\begin{figure} \\centerline{ \\includegraphics[width=\\singlecolsize\\textwidth]{f10.ps} } \\caption{Colour versus magnitude distribution for a near-IR sample.  The black contours and points represent potential galaxy targets.  The blue dashed lines show the limits imposed by our selection including the constraint $r_{\\rm model}<20.5$.  The green lines show $r=19.4$ and 19.8 limits. The red contours represent most of the additional targets not selected by the $r_{\\rm petro}$ limits. These contours extend below the solid green line because of differences between Petrosian and model magnitudes.} \\label{fig:color-bias} \\end{figure}  \\subsection{Masking} \\label{sec:masking}  In order to avoid targeting galaxies with bad photometry because they are near bright stars or satellite trails, an explicit mask was constructed.  The bright-stars mask was based on stars down to $V<12$ in the Tycho~2, Tycho~1 and Hipparcos catalogues. For each star, a scattered-light radius ($\\rscat$) was estimated based on the circular region over which the star flux per pixel is greater than 5 times the sky noise level.  For each potential target, a mask parameter was defined as follows \\begin{equation} \\begin{array}{l} \\mbox{{\\small MASK\\_IC\\_12}} = 1  \\\\ \\mbox{{\\small MASK\\_IC\\_12}} = \\rscat/d  \\\\ \\mbox{{\\small MASK\\_IC\\_12}} = 0 \\end{array} \\mbox{~for~} \\begin{array}{l} d \\le \\rscat  \\\\ \\rscat < d  \\le 5 \\rscat \\\\ d > 5 \\rscat \\end{array} \\end{equation} where $d$ is the distance to a $V<12$ star with radius $\\rscat$.  In other words, the {\\small MASK\\_IC\\_12} value decreases from unity when $d \\le \\rscat$ to 0.2 when $d = 5 \\rscat$.  A similar mask parameter {\\small MASK\\_IC\\_10} was defined using only $V<10$ stars. In addition, objects within an SDSS-database mask for holes, satellite trails and bleeding pixels had these mask values set to unity.  After testing, we chose to select only objects with {\\small MASK\\_IC\\_10} $< 0.5$ and {\\small MASK\\_IC\\_12} $< 0.8$.  The largest masked areas are shown in Fig.~\\ref{fig:ukidss-coverage}.  These are between 0.01 and $0.07\\,\\sqdeg$ each.  Most of the separate masked areas are significantly smaller ($<0.001\\,\\sqdeg$ or $<1'$ in radius).  Overall, the total masked area is about $1.0\\,\\sqdeg$ and the unmasked area of the survey is estimated to be $143.0\\,\\sqdeg$.  The mask was insufficient to remove all or nearly all objects with bad photometry.  Therefore, as per SDSS selection, objects were selected to be \\textsc{not satur} from the \\textsc{flags} column in the \\textsc{PhotoObj} table. This basically excludes deblends of bright stars but will also reject galaxies that are blended with saturated stars.  These however are likely to have bad photometry and falsely bright magnitudes.  The stars causing this saturation, not accounted for by the Tycho mask, are probably around $V\\sim13$.  The saturated-flag masking is not ideal.  This is particularly the case for large nearby galaxies for which the angular size of the galaxy is a significant factor in determining the excluded sky area.  In other words, the probability of a large galaxy having \\textsc{satur} set depends primarily on its size rather than the area of the diffracted and scattered light around stars.  To increase the completeness of the input catalogue for large galaxies, exceptions for the mask and not-saturated criteria were made for galaxies from the Uppsala General Catalog (UGC; \\citealt{cotton99}) and Updated Zwicky Catalog (UZC; \\citealt{falco99}).  In addition, exceptions to the not-saturated criteria were made for a selection of visually inspected galaxies that have \\textsc{not satur\\_center}. There are only 86 objects with an exception flag set (selected as part of the visual classification process described in \\S~\\ref{sec:vis-class}).  In summary, the criteria for including objects is given by: \\begin{equation} \\begin{array}{l} (\\mbox{\\textsc{mask\\_ic\\_10}} < 0.5 \\mbox{~~\\textsc{and}~~} \\mbox{\\textsc{mask\\_ic\\_12}} < 0.8 \\mbox{~~\\textsc{and}}  \\\\ \\mbox{\\textsc{not satur}}  )  \\mbox{~~\\textsc{or}~~} \\mbox{the exception flag is set.} \\end{array} \\label{eqn:masking} \\end{equation}  \\subsection{Surface brightness limits} \\label{sec:sb-limits}  In addition to the implicit surface brightness (SB) limits from star-galaxy separation and detection, an explicit SB limit was applied given by \\begin{equation} 15.0 < \\effsb < 26.0 \\label{eqn:sb-limits} \\end{equation} where $\\effsb$ is the effective SB in ${\\rm mag\\,arcsec}^{-2}$ within the 50\\% light radius in the $r$-band (eq.~5 of \\citealt{strauss02}). Anything of lower SB is very likely to be an artifact, and anything of higher SB is a star.  Figure~\\ref{fig:sb-sep} shows the distribution of objects in $r_{\\rm fibre}$ versus $\\effsb$ for GAMA main-survey targets. The lower limit of 15.0 does remove some objects, probably stars, not rejected by the masking or star-galaxy separation criteria (Eq.~\\ref{eqn:star-gal-summary}).  The limit for $\\effsb$ of $26.0$ is 1.5~magnitudes deeper than the SDSS MGS cut, and is the point at which most of the objects are clearly artifacts.  Note that the SDSS photometric pipeline is not complete for $\\effsb > 23$ (figs.~2--3 of \\citealt{blanton05}).  Additional low-SB candidates could be recovered by searching coadded $g$, $r$ and $i$ images \\citep{kniazev04}. Nevertheless without deeper imaging, the data will remain incomplete at low SB well before our explicit limit.  \\begin{figure*} \\includegraphics[width=\\middlecolsize\\textwidth]{f11.ps} \\caption{Bivariate distribution of $r_{\\rm fibre}$ versus $\\effsb$.  The black contours and points represent objects that are not masked and pass star-galaxy separation.  The grey lines outline the selection limits: $15.0<\\effsb<26.0$ is the restriction for the science catalogue; while objects with fibre magnitudes fainter than 22.5 or brighter than 17.0 are not included in the AAOmega observation schedule.  The green dots represent objects with \\visclass=1; the red crosses \\visclass=2, and the pink crosses \\visclass=3. The small red circles at $r_{\\rm fibre}<17$ are probably stars based on a stricter star-galaxy separation criteria (\\S~\\ref{sec:bright-galaxies}). The blue dash-dotted (dashed) line corresponds to 50\\% (30\\%) redshift success rate for objects on or near the line.} \\label{fig:sb-sep} \\end{figure*}  In addition to the explicit SB limits given in Eq.~\\ref{eqn:sb-limits}, which we use to reject objects from our science catalogue, we include a restriction on the fibre magnitudes: \\begin{equation} 17.0 < r_{\\rm fibre} < 22.5 \\label{eqn:fibre-limits} \\end{equation} for targets allocated to the AAOmega observation schedule.  This is a practical restriction, with a bright limit to avoid significant crosstalk in the spectrograph and a faint limit because the redshift success is very low. Selected fibre bright targets without a known redshift will be observed with a 2m-class telescope, and, in principle, selected fibre faint targets will be observed with an 8m-class telescope. At the bright end, a more restrictive cut on star-galaxy separation is also justified (see later in \\S~\\ref{sec:bright-galaxies}).  \\subsection{Visual classification} \\label{sec:vis-class}  Sources with, for example, $\\effsb > 23$ have a high probability of being artifacts, deblends of stars, or the outer parts of galaxies.  One of us (J.\\ Liske) has written code to facilitate the visual classification of such sources. A \\visclass\\ variable, initially with zero value, could be changed to the following for each source on inspection: \\begin{itemize} \\item [1] possibly a target, \\item [2] not a target (no evidence of galaxy light), \\item [3] not a target (not the main part of a galaxy). \\end{itemize}  First, sources with the following flags all equal to zero, \\textsc{edge, blended, child, maybe\\_cr, maybe\\_eghost}, were assumed to be good, essentially isolated, and not included in any testing (\\visclass\\ set to 255). About 50\\% of targets satisfy these criteria.  From the remaining objects, sources were selected for visual classification if any of the following conditions applied: $\\effsb > 23$, $r_{\\rm fibre} > 21$, $r_{\\rm fibre} < 17$, \\textsc{mask\\_ic\\_12} $> 0.2$, $r_{\\rm model} < 15.5$, $r_{\\rm petro} < 15.5$, $r_{\\rm fibre} < r_{\\rm model}$, $r_{\\rm fibre} < r_{\\rm petro}$, near UGC galaxy, within $3''$ of another target, Petrosian radius $>10''$.  These indicate that the object could be the result of deblending of a large galaxy, artifact or bright star, e.g., diffraction spikes.  In addition to the above criteria, other objects were included in the above process.  Objects with the same \\textsc{parentid} as an already classified \\visclass\\ = 3 object were selected. (The above selection was not developed in one go and there have been several iterations.) Finally objects, with the same \\textsc{parentid}, that are the brightest and nearest to any object to be tested were included. Objects that could be part of the same galaxy were viewed together where possible. One had to be certain to classify objects as 3 only if the main part was identified as a target.  The above selection produced a sample of about 12\\,500 objects for visual classification, by six observers. Every selected object was classified by three different observers.  Of the selected potential main-survey targets (Fig.~\\ref{fig:sb-sep}), \\visclass=1 was set in 92\\% of cases, \\visclass=2 in 5\\% of cases, and \\visclass=3 in 3\\% of cases, based on agreement between two or all three classifiers, 9\\% and 90\\% of cases, respectively.  Some of the ambiguous cases were double checked, and a single-observer classification was selected in 1\\% of cases.  Objects with values of 2 or 3 were removed from the schedule of AAT observations, i.e., targets must satisfy \\begin{equation} \\mbox{\\visclass} \\neq 2 \\mbox{~~\\textsc{and}~~} \\mbox{\\visclass} \\neq 3 \\: . \\label{eqn:vis-class} \\end{equation} In addition, the $\\visclass=3$ objects can be used to improve the photometry of some large galaxies by coadding in the flux of the galaxy parts (or the `parent' photometry can be used).  \\subsection{Number of targets} \\label{sec:number-targs}  The total number of objects that are within the GAMA regions (\\S~\\ref{sec:sdss}), main-survey magnitude limits (Eq.~\\ref{eqn:mag-limits}) and $\\dsg > 0.05$, is 143\\,728. Applying the stricter star-galaxy separation (Eq.~\\ref{eqn:star-gal-summary}) reduces the sample to 132\\,073.  Removing objects by masking (Eq.~\\ref{eqn:masking}), the SB limits (Eq.~\\ref{eqn:sb-limits}) and visual checking (Eq.~\\ref{eqn:vis-class}), reduces the sample to 120\\,038. Of these, 825 were not included in the AAOmega observation schedule because they do not satisfy the fibre magnitude limits (Eq.~\\ref{eqn:fibre-limits}).  A more restictive star-galaxy separation can be applied for brighter targets (discussed later and given in Eq.~\\ref{eqn:star-gal-special}) that reduces the sample to {\\bf 119\\,852}. This is considered to be the main-survey sample.  Note these numbers apply to AAT observations in 2009, the numbers may change slightly with addition of complete $J$-$K$ UKIDSS.  Separating the main survey into $r$, $z$ and $K$ limited samples, the numbers are 114\\,520, 61\\,418 and 57\\,657, respectively.  \\newtext{For $r_{\\rm petro}$ selected samples to 19.0, 19.4 and 19.8, the sample sizes are 60\\,407, 96\\,386 and 150\\,810, respectively.  The latter is the $r$-selected main survey plus F2 additional targets, which are described in the following section.}  \\subsection{Additional targets} \\label{sec:filler-targets}  In order to assess the spectro-photometry of the AAOmega spectra, three or four stars, classified as \\textsc{redden\\_std} or \\textsc{spectrophoto\\_std} by SDSS, were observed in each configuration. These also had a bright fibre magnitude limit of 17 as per the main-survey targets.  The aim is to obtain high completeness (99\\%), at least in terms of spectra obtained and ideally in terms of confirmed redshifts, for the main survey. This is set to reduce systematic uncertainties in GAMA's position dependent science cases, and is possible because a given patch of sky is potentially observed by $\\sim5$--10 2dF tiles depending on the local density of targets (see \\citealt{robotham09} for a description of the tiling strategy).  Given this requirement, targeting becomes increasingly inefficient as the survey progresses (fewer targets without a redshift per tile). Filler targets were introduced to provide useful redshifts outside the main survey, and thus, maximise fibre usage.  These have no high-level requirement on completeness. The filler selections are given by: (F1) objects with detection in the Faint Images of the Radio Sky at Twenty-cm (FIRST) survey and matched to SDSS with $i_{\\rm model} < 20.5$ including unresolved sources; (F2) $19.4 < r_{\\rm petro} < 19.8$ galaxy targets in G09 and G15, aiming for equal depth with G12; and (F3) $g_{\\rm model}<20.6$ or $r_{\\rm model}<19.8$ or $i_{\\rm model}<19.4$ in G12, investigating variation in magnitude-type and wavelength on selection. In total, there are about 50\\,000 filler targets.   \\label{sec:summary}  The GAMA survey is designed to be a highly complete redshift survey with a target density several times that of SDSS. The survey covers three $48\\,\\sqdeg$ regions near the celestial equator centred on 9\\,h, 12\\,h and 14.5\\,h (Fig.~\\ref{fig:sdss-stripes}).  The input catalogue is drawn from the SDSS and UKIDSS.  The main-survey limits are $r_{\\rm petro}<19.4$, $z_{\\rm model}<18.2$ and $K_{\\rm AB,auto}<17.6$ ($K<15.7$) across all the regions, and $r_{\\rm petro}<19.8$ over the G12 region (Eq.~\\ref{eqn:mag-limits}).  This corresponds to a main survey of 119\\,852 targets.  The near-IR selections have a joint constraint with $r_{\\rm model}<20.5$, which has minimal impact on the use of the near-IR selections (Figs.~\\ref{fig:color-bias} \\&~\\ref{fig:obs-color-z}). The GAMA survey lies between that of the SDSS-MGS $r<17.8$ and VVDS-wide $I_{\\rm AB}<22.5$ magnitude-limited samples in the depth-area plane (Fig.~\\ref{fig:compare-grs}).  In terms of $K$-band selection (Figs.~\\ref{fig:ukidss-coverage} \\&~\\ref{fig:K-counts}), GAMA covers an area $\\sim15$ times that of the similar-depth Hawaii+AAO $K<15$ survey.  In order to be highly complete at the high-SB end of the galaxy distribution, an intensity profile parameter (Eq.~\\ref{eqn:sdss-star-gal-sep}) and a colour-colour parameter (Eq.~\\ref{eqn:ukidss-sdss-star-gal-sep}) are used jointly for star-galaxy separation.  The $\\dsgjk$ parameter makes use of $J-K$ and $g-i$ colours.  Either parameter works reasonably well in separating stars and galaxies (Figs.~\\ref{fig:sg-sep-histo}--\\ref{fig:test-auto-jk}). A joint selection (Eq.~\\ref{eqn:star-gal-summary}) increases the completeness while stellar contamination in the sample remains at less than 2\\%.  Judging by the joint distribution of confirmed galaxies in these parameters (Fig.~\\ref{fig:results-sg-sep}), the completeness is high because the bivariate density drops significantly prior to the limit of our selection. This is particularly important when considering the size evolution of galaxies (Fig.~\\ref{fig:size-z}). The incompleteness at the low-SB end is significant, both in source detection and redshift success rate, which is about 50\\% at $r_{\\rm fibre}=21.5$ (Fig.~\\ref{fig:sb-sep}).  Some improvement over the SDSS MGS is made by visually checking low-SB targets ($\\effsb > 23$), rather than using automatic checks, by increased redshift success rate, and by eventually including further integrations of sources with failed redshifts.  The GAMA survey has completed two out of a three-year time allocation for spectroscopy with AAOmega on the AAT. To date, 100\\,012 redshifts have been confirmed for the main survey, including 80\\,944 from AAOmega.  Of these, 98.5 per cent are extragalactic.  \\newtext{The completeness is 96\\% for $r_{\\rm petro}<19.0$, 74\\% for the fainter $r$-band selection, and $\\sim39$\\% for the remaining near-IR selection (Table~\\ref{tab:gama-main-spec}).  The completeness at $r>19$ will be significantly improved in the third year of spectroscopic observations.}  We expect that this galaxy redshift survey will form a core of a fundamental database for many studies in extragalactic astronomy."
    },
    "0910/0910.0858.txt": {
        "abstract": "We present a maximum-likelihood analysis for estimating the angular distribution of power in an anisotropic stochastic gravitational-wave background using ground-based laser interferometers. The standard isotropic and gravitational-wave radiometer searches (optimal for point sources) are recovered as special limiting cases. The angular distribution can be decomposed with respect to {\\em any} set of basis functions on the sky, and the single-baseline, cross-correlation analysis is easily extended to a network of three or more detectors---that is, to multiple baselines. A spherical harmonic decomposition, which provides maximum-likelihood estimates of the multipole moments of the gravitational-wave sky, is described in detail. We also discuss: (i) the covariance matrix of the estimators and its relationship to the detector response of a network of interferometers, (ii) a singular-value decomposition method for regularizing the deconvolution of the detector response from the measured sky map, (iii) the expected increase in sensitivity obtained by including multiple baselines, and (iv) the numerical results of this method when applied to simulated data consisting of both point-like and diffuse sources. Comparisions between this general method and the standard isotropic and radiometer searches are given throughout, to make contact with the existing literature on stochastic background searches. ",
        "introduction": "\\label{s:intro}  Data from the laser interferometric gravitational-wave detectors LIGO \\cite{LIGO, Barish:1999,LIGO_eprint,detector}, Virgo \\cite{Virgo, Bradaschia:1990}, and GEO \\cite{GEO-600, Wilke:2004} are currently being analysed for the presence of gravitational waves from a variety of sources. These include signals from inspiraling and coalescing compact binaries (for example, neutron stars and/or stellar mass black holes)~\\cite{bin-insp-S3S4,lowmass_S5,ringdowns_S4,lowmass_S5B}, continuous gravitational waves from quasi-periodic sources such as pulsars~\\cite{SGR,SGR1806,radiopulsars,allsky}, and bursts of gravitational radiation associated with gamma-ray bursts~\\cite{GRB070201,39GRBs,S4GRB}, core-collapse supernovae, or other violent events~\\cite{S3bursts}. In addition, searches are ongoing for the presence of a background of {\\em stochastic} gravitational radiation of either astrophysical or cosmological origin~\\cite{radiometer,LIGO-ALLEGRO,S4Isotropic}, whose detection might provide insights about the very early universe \\cite{Maggiore:2000}, well before the production of the cosmic microwave background \\cite{Kolb:1999}.  Although no direct detections of gravitational waves have been made to date, the most recent data taken are of unprecented sensitivity~\\cite{stoch-S5,SGR,GRB070201}, leading to upper limits on gravitational-wave strengths that are competitive with or surpass those from electromagnetic or particle physics observations. Of particular note is the upper limit on the strength of a gravitational-wave signal from the Crab pulsar \\cite{LIGO-crab}, which is a factor of 1.6 lower than the corresponding limit inferred from electromagnetic pulsar spin-down observations~\\cite{palomba}. Also, the current direct limit on the strength of an isotropic stochastic gravitational-wave background at $\\unit[100]{Hz}$ $\\Omega_{\\text gw}<6.9\\times10^{-6}$~\\cite{stoch-S5} (at 95\\% confidence) has surpassed bounds set by considerations from Big Bang Nucleosynthesis~\\cite{Maggiore:2000} and from the microwave background~\\cite{CMB-limit}.  In this paper, we describe an analysis method that estimates the angular distribution of power in an {\\em anisotropic} stochastic gravitational-wave background. This method includes both the standard isotropic \\cite{S1Isotropic,S3Isotropic,S4Isotropic} and gravitational-wave radiometer searches \\cite{mitra-et-al,radiometer} (optimal for point sources) as special limiting cases. (For our purposes {\\em anisotropic} is taken to mean {\\em not necessarily isotropic}.) Similar to the radiometer technique, the method presented here looks for modulations in the gravitational-wave signal induced by the Earth's rotational motion relative to an anisotropic background. The method provides maximum-likelihood estimates of the angular distribution of gravitational-wave power ${\\cal P}(\\hat\\Omega)=\\sum_\\alpha {\\cal P}_\\alpha{\\bf e}_\\alpha(\\hat\\Omega)$, decomposed with respect to some set of basis functions on the sky. By choosing a pixel basis %${\\cal P}_{\\hat \\Omega}={\\cal P}(\\hat\\Omega)$, ${\\bf e}_{\\hat\\Omega'}(\\hat\\Omega)=\\delta(\\hat\\Omega,\\hat\\Omega')$, we recover the results of the radiometer method discussed in~\\cite{mitra-et-al,radiometer}. By choosing the spherical harmonics basis $Y_{lm}(\\hat\\Omega)$ defined with respect to the Earth's rotational axis, we obtain maximum-likelihood estimates of the {\\em multipole moments} ${\\cal P}_{lm}$ of the gravitational-wave sky. This basis is particularly convenient as the standard isotropic analysis corresponds to simply restricting attention to the monopole moment ${\\cal P}_{00}$, while the point-source radiometer results are well-approximated by choosing a sufficiently large value of $l_{\\rm max}$ ($l_{\\rm max}\\sim30$), appropriate for the diffraction-limited beam pattern at $f\\sim 1$~kHz. In addition, the use of spherical harmonics simplifies the problem of removing the `smearing' effects of the beam pattern from the measured sky map (that is, deconvolution of the {\\em dirty map}), given the smaller number of elements and symmetries of the beam pattern matrix with respect to the $lm$ indices. The problem of deconvolving a cross-correlated gravitational-wave signal from the interferometers' beam pattern in the spherical harmonic basis has been described in~\\cite{Cornish:2001hg}. We address this problem in detail in section~\\ref{s:data_analysis}. We further note that the spherical harmonic basis is useful for the efficient analysis of cross-correlated data in a variety of applications including searches for transient gravitational-wave sources~\\cite{kipp}.  Regardless of basis, the method described here is easily extended to work with a network or three or more detectors with uncorrelated detector noise, by simply adding the individual baseline beam patterns and dirty maps before deconvolution. A multi-baseline analysis improves the overall sensitivity of the search by reducing the variances of the individual estimators, and provides a natural way of regularising the deconvolution of the dirty map; the beam pattern matrix has fewer null (or nearly null) directions for multiple baselines and is thus more stable during inversion.  The structure of the rest of the paper is the following: In section~\\ref{s:sgwb}, we briefly review the statistical properties of an anisotropic background, and show how a generalized overlap reduction function arises in a cross-correlation search for such a background. In section~\\ref{s:MLestimation} we derive the optimal estimators of the angular distribution of the gravitational-wave power, starting from the likelihood function for cross-correlated data. We explicitly construct the beam pattern matrix, and discuss its relation to the covariance matrix of the estimated ${\\cal P}_{\\alpha}$. Section \\ref{s:data_analysis} describes details of the data analysis implementation and issues related to deconvolution and regularisation. %Section \\ref{s:multiple_baselines} It also briefly describes how to extend %the results of the previous sections the analysis to a network of three or more detectors, and the expected increase in sensitivity from using multiple baselines. In section~\\ref{s:simulations} we present numerical results of the method applied to simulated data. We consider both point-like and diffuse-source injections, and compare the extracted and injected sky-maps. Finally, in section~\\ref{s:summary} we summarize our results. We also include three appendices: Appendices ~\\ref{s:Ylms} and \\ref{s:identities} contain definitions of the spherical harmonics and some useful identities relating different multipole moments, beam pattern matrix components, etc.; Appendix \\ref{s:DetStat} defines a related detection statistic that assumes a particular distribution of (normalized) angular distribution functions on the sky. ",
        "conclusions": "\\label{s:summary}  We have presented here a maximum-likelihood analysis method for estimating the angular distribution of power in an anisotropic stochastic gravitational-wave background. The basic idea was to cross-correlate data from a network of two or more gravitational-wave detectors, exploiting time-of-arrival differences and the diurnal modulation due to the Earth's rotation. We derived maximum-likelihood estimators for the angular distribution of gravitational-wave power ${\\cal P}(\\hat\\Omega)=\\sum_\\alpha {\\cal P}_\\alpha \\mathbf{e}_\\alpha(\\hat\\Omega)$, decomposed with respect to {\\em any} set of basis functions on the sky. We derived an expression for the beam pattern matrix $\\Gamma_{\\alpha\\beta}$ and discussed its relationship to the covariance matrix %$(\\Gamma^{-1})_{\\alpha\\beta}$ of the maximum-likelihood estimators $\\hat{\\cal P}_\\alpha$. We described how singular value decomposition can be used to regularize the inverse of $\\Gamma_{\\alpha\\beta}$, which was needed to remove the smearing effects of the beam pattern matrix on the measured (`dirty') sky maps $X_\\alpha$. We also explained how the single-baseline (two-detector) cross-correlation analysis can be extended to a network of three or more detectors, thereby increasing our sensitivity to detecting a signal. In this paper, we focused attention on a decomposition with respect to a basis of spherical harmonics $Y_{lm}(\\hat\\Omega)$, for which the maximum-likelihood estimators $\\hat{\\cal P}_{lm}$ represent the multipole moments of the gravitational-wave sky, and for which the standard isotropic %\\cite{S1Isotropic,S3Isotropic,S4Isotropic} and radiometer %\\cite{mitra-et-al, radiometer} searches are recovered as special limiting cases. Finally, we illustrated all these general results by analysing simulated data containing injected stochastic gravitational-wave backgrounds having different angular power distributions."
    },
    "0910/0910.1779_arXiv.txt": {
        "abstract": "The radio emitting X-ray binary GRS\\,1915+105 shows a wide variety of X-ray and radio states. We present a decade of monitoring observations, with the RXTE-ASM and the Ryle Telescope, in conjunction with high-resolution radio observations using MERLIN and the VLBA. Linear polarisation at 1.4 and 1.6~GHz has been spatially resolved in the radio jets, on a scale of $\\sim150$~mas and at flux densities of a few mJy. Depolarisation of the core occurs during radio flaring, associated with the ejection of relativistic knots of emission. We have identified the ejection at four epochs of X-ray flaring. Assuming no deceleration, proper motions of 16.5 to 27~mas per day have been observed, supporting the hypothesis of a varying angle to the line-of-sight per ejection, perhaps in a precessing jet. ",
        "introduction": "GRS\\,1915+105 has proved to be one of the best laboratories for the study of relativistic physical environments, due to its high-brightness, periodic outbursts of superluminal ejecta and relative proximity. Since the first detection in August 1992 with the WATCH instrument on board the X-ray telescope GRANAT~\\citep{1992IAUC.5590....2C}, GRS\\,1915+105 has demonstrated some of the most spectacular high-energy physics within the Galaxy. Shortly after its discovery, a radio counterpart was discovered by \\cite{1993IAUC.5773....2M} and later an infrared component was also identified \\citep{1993IAUC.5830....1M}. Its radio light-curve shows a highly variable and complex structure, with spatially resolved features that can be followed over many days.  A clear correlation between the X-ray, infrared and radio emission was quickly established by~\\cite{1994Natur.371...46M}, with rapid time variability in each band. Radio flaring was found to correspond to a rapid change in the hard X-rays and possible production of high-energy gamma rays. In March 1994, 20~cm VLA observations of GRS\\,1915+105 detected the first superluminal motion of a Galactic source \\citep{1994Natur.371...46M}. This major breakthrough provided the direct evidence of relativistic jets and an extreme physical environment within the Galaxy. The name ``microquasar'' was coined due to their obvious similarities with their extra-galactic counterparts, quasars.  The launch of the Rossi X-ray Timing Explorer (RXTE) satellite, in December 1995, signified the start of a long-term monitoring campaign of X-ray binaries. The All Sky Monitor (ASM) on board the RXTE has taken daily observations of GRS\\,1915+105 since its launch. \\cite{1996ApJ...473L.107G} detected unusual X-ray variability on time scales of under one second to days. The RXTE-ASM lightcurve was found to be both highly complex and structured, due to instabilities in the accretion disc. Detailed observations with the RXTE's Proportional Counter Array (PCA) instrument (\\citealt{1997ApJ...477L..41C}; \\citealt{1997ApJ...488L.109B}) identified different spectral states, including a low-hard state dominated by a power-law and a high-soft state with a strong disc-blackbody component. A transition between such states is believed to be associated with the ejection of superluminal plasmons \\citep{1999MNRAS.304..865F}.  \\begin{table*}  \\begin{center}\\begin{tabular}{|>{\\columncolor[rgb]{0.8,0.8,0.8}}l|ccccccc|} \\hline \\rowcolor[rgb]{0.8,0.8,0.8} Epoch (MJD) & Frequency/Array  & $I_{\\rm{total}}$ &  $LP_{\\rm{total}}$ &   Fraction & Noise & Spectral$^{\\ddagger}$ & Resolved? \\\\ \\rowcolor[rgb]{0.8,0.8,0.8}      & (GHz) &  (mJy)       &   (mJy) &   (\\%)            & ($\\mu$Jy/beam)    &  Index  &      \\\\  \\hline $^{\\star}$1997 October 31 (50752) -core & 2.27 / VLBA &  $105\\pm2$ &    $\\le0.1^{\\dagger}$ &  $\\le0.5^{\\dagger}$ & 430 & \\multirow{4}{*}{-} & \\multirow{4}{*}{YES} \\\\ $^{\\star}$1997 October 31 (50752) -jet & 2.27 / VLBA & $226\\pm2$ &  $5.6\\pm0.3$ & 2.5 & 430 & & \\\\ $^{\\star}$1997 October 31 (50752) -core & 8.3 / VLBA  & $33\\pm1$ &    $\\le0.1^{\\dagger}$& $\\le0.5^{\\dagger}$ & 460 & &\\\\ $^{\\star}$1997 October 31 (50752) -jet & 8.3 / VLBA  & $65\\pm3$ & $4.7\\pm0.3$& 7.2  & 460 & &\\\\ \\hline 1998 June 9 (50973) & 1.6 / MERLIN &  $145.1\\pm0.4$ &  $9.5\\pm0.3$&  $6.9\\pm0.2$ & 300  & -0.3 & NO\\\\ \\hline 1999 April 2 (51270) & 1.6 / MERLIN &  238.4$\\pm0.3$ &  $3.5\\pm0.3$& $1.5\\pm0.2$ & 530  & -0.6& NO\\\\ \\hline 1999 November 18 (51500) & 1.6 / MERLIN &  221.3$\\pm0.8$ &  $6.9\\pm0.3$ & $3.2\\pm0.2$  & 860 & -0.4& NO\\\\ \\hline 1999 November 22 (51504) & 1.6 / MERLIN &  37.1$\\pm0.7$ &  $\\le0.1^{\\dagger}$ & $\\le0.4^{\\dagger}$ & 300 & -0.2& NO\\\\ \\hline 1999 December 28 (51540) & 1.6 / MERLIN &  89.1$\\pm0.7$ & $12.7\\pm1.6$  & $14.1\\pm0.2$ & 400 & -0.4& NO\\\\ \\hline 2003 March 6 (52704) & 1.6 / MERLIN &  $31.3\\pm0.3$ &  $1.6\\pm0.3$ & $5.2\\pm1.0$ & 240 & -0.4& NO\\\\ \\hline 2003 March 25 (52723) & 1.6 / MERLIN & $125.8\\pm0.1$ &  $7.1\\pm0.2$ & $5.6\\pm0.2$  & 140 & 0.2& NO\\\\ \\hline 2003 April 18 (52747) -core & 1.6 / MERLIN &  $135.4\\pm0.4$ &  $3.5\\pm0.2$ & $2.6\\pm0.2$  & 270 & \\multirow{2}{*}{0.0}& \\multirow{2}{*}{YES}\\\\ 2003 April 18 (52747) -jet & 1.6 / MERLIN &  $20.0\\pm0.3$ &  $1.3\\pm0.1$ & $6.5\\pm0.5$  & 270 & &\\\\ \\hline 2003 May 09 (52768) & 1.6 / MERLIN &  $43.0\\pm0.1$ &  $2.7\\pm0.1$ & $6.3\\pm0.2$  & 160 & 0.2&NO\\\\ \\hline 2003 June 15 (52805) -core & 1.6 / MERLIN &  $41.4\\pm0.1$ &  $\\le0.1^{\\dagger}$  & $\\le0.2^{\\dagger}$ & 390 & \\multirow{2}{*}{-0.4}& \\multirow{2}{*}{YES}\\\\ 2003 June 15 (52805) -jet & 1.6 / MERLIN &  $67.5\\pm0.1$ &  $7.6\\pm0.4$  & $11.9\\pm0.6$ & 390 & & \\\\ \\hline 2006 December 24 (54093) & 1.4 / MERLIN &  $145.0\\pm0.5$ &  $7.7\\pm0.6$  & $5.3\\pm0.4$ & 380 & \\multirow{2}{*}{-0.1}& \\multirow{2}{*}{NO}\\\\ 2006 December 24 (54093) & 1.6 / MERLIN &  $156.9\\pm1.2$ &  $11.5\\pm0.5$ &  $7.3\\pm0.5$ & 380 && \\\\ \\hline 2006 December 27 (54096) & 1.4 / MERLIN &  $20.7\\pm0.5$ &  $\\le0.1^{\\dagger}$  & $\\le0.5^{\\dagger}$ & 230 & \\multirow{2}{*}{-0.4} & \\multirow{2}{*}{NO}\\\\ 2006 December 27 (54096) & 1.6 / MERLIN &  $27.0\\pm0.5$ &  $\\le0.1^{\\dagger}$  & $\\le0.5^{\\dagger}$ & 230 & &\\\\ \\hline 2006 December 28 (54097) & 1.4 / MERLIN &  $25.8\\pm0.4$ &  $\\le0.1^{\\dagger}$  & $\\le0.4^{\\dagger}$ & 210 & \\multirow{2}{*}{-}& \\multirow{2}{*}{NO}\\\\ 2006 December 28 (54097) & 1.6 / MERLIN &  $26.2\\pm0.4$ &  $\\le0.1^{\\dagger}$  & $\\le0.4^{\\dagger}$ & 210 & &\\\\ \\hline 2007 January 04 (54104) & 1.4 / MERLIN &  $21.0\\pm0.5$ &  $\\le0.1^{\\dagger}$ & $\\le0.5^{\\dagger}$ & 260 & \\multirow{2}{*}{-0.3}& \\multirow{2}{*}{NO}\\\\ 2007 January 04 (54104) & 1.6 / MERLIN &  $22.2\\pm0.5$ &  $\\le0.1^{\\dagger}$  & $\\le0.5^{\\dagger}$ & 260 & &\\\\ \\hline \\end{tabular}\\end{center} \\caption{\\label{table:journal}Journal of observations and results: Date of observation, observed frequency and array, total flux, total linearly polarised flux, fraction of total flux that is linearly polarised, systematic noise, spectral index and detection of resolved structure. [Notes: $^{\\star}$ The observations were first published by \\citealt{2000ApJ...543..373D}; $^{\\dagger}$ Estimates are based on a upper limit of detection; $^{\\ddagger}$ The spectral index is derived from the total MERLIN flux (i.e. core + jet) and the 15 GHz RT data, where $S \\propto \\nu^\\alpha$.]} \\end{table*}  Radio polarisation was first detected with MERLIN in GRS\\,1915+105 by \\cite{1999MNRAS.304..865F} in the first four epochs (i.e. first four days) of the October 1997 flare at a frequency of 5~GHz. A clear asymmetry in the location of the polarised emission was found, showing the strongest detection of linearly polarised emission within the approaching jet, weaker in the receding jet and no detection ($<3\\%$ of the total flux density) within the core. The polarisation position angle (PA) was found to rotate by at least $75^\\circ$ over the four days, leading to the suggestion that the changes in polarisation could be due to Faraday effects (i.e. changes in the Faraday depth as the plasmons move away from the core), implying a change in rotation measure of $\\Delta RM>300$~rad~m$^{-2}$.  This paper describes observations taken over a decade with MERLIN at 18 and 21~cm. The observations were either in conjunction with INTEGRAL observations as part of the Galactic Plane Survey~\\citep{2004ApJ...607L..33B,2004ESASP.552..321F}, or were triggered by flare events, found by other radio telescopes. Polarisation behaviour, structural variations and the relationship of activity in the radio regime to the X-ray behaviour were investigated. ",
        "conclusions": "GRS\\,1915+105 has been simultaneously observed in the radio and X-rays in various different states over a ten year period. This comparative study between the X-ray and radio variability shows, for the first time, both long term ($\\sim$~weeks/months) and short lived (intra-day) variations. Daily monitoring shows periods of extended activity and X-ray spectral states that are related to the ejection of material. Radio activity over this period has shown to originate from the central 150 mas core, unless a short X-ray flare is observed that marks the ejection of a superluminal plasmon.  The proper motion of ejecta has been calculated by identifying the time of ejection and exhibit no deceleration. Measured velocities were in agreement with previous observations, and show a significant variation in velocity. Polarisation measurements show strong linear polarisation variation. Depolarisation in the core of the XRB is found during the ejection of a plasmon and the multi-wavelength observations show an increase of linear polarisation with frequency, suggesting internal Faraday rotation effects.  Observing GRS\\,1915+105 over a long period of time has shown the relativistic out-flow in different forms. If the accretion conditions permit, collimated and steady jets persist over many weeks. Different apparent  velocities of superluminal knots suggest either a precession of the jet angle or variation in their formation. Finally, the nature of the short-lived radio flares remains unknown; for example, it is difficult to explain why these bright, but short-lived events do not produce extended structure, in contrast to a flare during a state transition."
    },
    "0910/0910.4640_arXiv.txt": {
        "abstract": "We processed the data about radial velocities and HI linewidths for $1678$ flat edge-on spirals from the Revised Flat Galaxy Catalogue. We obtained the parameters of the multipole components of large-scale velocity field of collective non-Hubble galaxy motion as well as the parameters of the generalized Tully-Fisher relationship in the ``HI line width -- linear diameter'' version. All the calculations were performed independently in the framework of three models, where the multipole decomposition of the galaxy velocity field was limited to a dipole, quadrupole and octopole terms respectively. We showed that both the quadrupole and the octopole components are statistically significant.  On the basis of the compiled list of peculiar velocities of $1623$ galaxies we obtained the estimations of cosmological parameters $\\Omega_m$ and $\\sigma_8$. This estimation is obtained in both graphical form and as a constraint of the value $S_8=(\\Omega_m/0.3)^{0.35} \\sigma_8 = 0.91 \\pm 0.05$. ",
        "introduction": "\\label{s:Introduction}  The distribution of matter, including dark matter, in the Universe is inhomogeneous on the scales of about $100\\,Mpc$. This is manifested, for example, in the existence of superclusters and voids. A galaxy, besides the cosmological expansion, is also attracted to the regions with greater density. As a result, the galaxies are involved in a non-Hubble large-scale collective motion on the background of Hubble expansion. In a more sophisticated way this can be considered as a development of initial density and velocity fluctuations in the early Universe due to the gravitational instability.  Investigation of such a motion is important at least for two reasons. First of all, it allows to plot the distribution of matter, including dark matter, in the surrounding region of the Universe, which is the main goal of cosmography, and to compare this distribution with the distribution of luminous matter. The second reason is that the parameters of this motion are linked with certain cosmological parameters, so we can obtain new independent estimations of these parameters. Of course, the accuracy of such estimation will not be very high, but the agreement of different estimations of cosmological parameters can support the correctness of the standard cosmological model.  In 1989 \\citeauthor{ref:K89} proposed to use flat edge-on spiral galaxies as a tool for studying their large-scale collective motion. They are good in this role for the following reasons: \\begin{enumerate} \\item{The linear diameter is strongly correlated with the HI linewidth for thin bulgeless galaxies. This allows to determine distances without photometric data.} \\item{Flat galaxies can be easily identified by their axes ratio.} \\item{Flat galaxies have a nearly 100 per cent HI detection rate.} \\item{Galaxies in clusters are usually not flat due to interaction with neighbours. This means that flat galaxies avoid clustering and do not interact with the intergalactic gas in clusters.} \\item{Peculiar velocities of isolated flat galaxies are not perturbed by neighbours.} \\end{enumerate}  To implement this idea the Flat Galaxy Catalogue \\citep[FGC,][]{ref:FGC} was created. It contained data about $N=4455$ galaxies, which satisfied the conditions $a_b/b_b\\ge7$ and $a_b>0.6'$. Here $a_b$ and $b_b$ are the major and minor axial diameters directly measured on POSS-I and ESO/SERC plates. In accordance with the original photographic material, the Catalogue consists of two parts: FGC ($N=2573$) and its southern extension, FGCE ($N=1882$). The first part is based on the POSS-I and covers the sky region with declinations between $-20\\deg$ and $+90\\deg$. The second one is based on the ESO/SERC and covers the rest of the sky area up to $\\delta=-90\\deg$.  After thorough studies of the catalogue's properties, both these parts were joined. The angular diameters from the FGCE were converted to the POSS-I system of the FGC, which appeared to be close to the $a_{25}$ system. This system, where galaxy size is taken at $B=25\\,mag/\\square''$ isophotal level, was used, in particular, by \\citet{ref:deVac}. Some FGCE galaxies, which did not satisfy the condition $a>0.6'$ after conversion, were removed from the resulting Revised Flat Galaxy Catalogue \\citep[RFGC,][]{ref:RFGC}. It contained data about $N=4236$ galaxies including the information on the following parameters: Right Ascension and Declination for the epochs J2000.0 and B1950.0, galactic longitude and latitude, major and minor blue and red diameters in arcminutes in the POSS-I diameter system, morphological type of the spiral galaxies according to the Hubble classification, index of the mean surface brightness (I -- high, IV -- very low) and some other parameters, which are not used in this article. More detailed description of the catalogue can be found in the paper \\citep{ref:RFGC}.  The original goal of this catalogue was to estimate the distance to galaxies according to the Tully-Fisher relationship in the ``HI line width -- linear diameter'' version without using their redshifts. The difference between the velocity $V$ derived from the redshift and the Hubble velocity $R=Hr$ corresponding to the distance $r$ estimated by Tully-Fisher relationship is called a peculiar velocity $V_{pec}=cz-Hr$. We can use such a simple form of the Hubble's law because our sample has $z<0.1$.  There are some important things to take into account about peculiar velocities. The redshift includes not only the velocity of the galaxy, but also the velocities of our Galaxy, Solar System and the Earth. Thus, to eliminate these factors, all velocities were reduced to the frame, where CMB is isotropic. Naturally, the redshift gives us only the radial component of the velocity and the tangential components cannot be measured. Additionally, Tully-Fisher relationship is statistical and thus has a certain error. Thus, we can only estimate the peculiar velocity for each galaxy, sometimes with a significant error.  However, we believe that there is a large-scale velocity field. We consider the individual galaxies as test particles in the velocity field of large-scale collective motion. Thus, having data about the peculiar velocities of a large number of galaxies, we can restore the radial component of the velocity field. Using some additional assumptions, like the potentiality of the flux, it is possible to restore the 3D velocity field. For this reason we need ample samples of peculiar velocities, preferably uniformly covering the celestial sphere.  These conditions are satisfied by the RFGC catalogue, which covers the whole celestial sphere in both hemispheres with the natural exclusion of the Milky Way region. However, not all of the RFGC galaxies have data required to estimate the distance to them using the Tully-Fisher relationship. Nevertheless, the sample of galaxies having such data is also quite uniform, as shown on Fig. \\ref{fig:1}.  \\begin{figure*}[tb] \\includegraphics[width=\\textwidth]{Parnovsky_fig1.eps} \\caption{Distribution of galaxies over the celestial sphere in galactic coordinates} \\label{fig:1} \\end{figure*}  Naturally, RFGC is not the only sample that can be used for this purpose. Possible alternatives include SBF \\citep{ref:SBF97,ref:SBF00}, ENEAR \\citep{ref:ENEAR}, SFI \\citep{ref:SFIa,ref:SFIb}, SFI++ \\citep{ref:SFIpp}, EFAR \\citep{ref:EFAR96,ref:EFAR01}, SMAC \\citep{ref:SMAC99,ref:SMAC04}, 2MFGC \\citep{ref:2MFGC}. However, in this article we use only the RFGC catalogue.  To apply the Tully-Fisher relationship the data from the catalogue is not sufficient; we also need to know the HI linewidth $W$ (in this article we take it at $50$ per cent level), or the gas rotation velocity $V_{rot}$ obtained from optical observations. Additionally, we need redshift data. The number of galaxies with such data increases constantly.  Since FGC and RFGC were assembled, a number of articles were published dealing with collective motions of RFGC galaxies. Very preliminary results were published by \\citet{ref:K95}. The parameters of the bulk motion were calculated by \\citet{ref:K00}. In the paper \\\\ \\citep{ref:Par01} not only new data were added, but also the models featuring the quadrupole and octopole components of the velocity field were proposed. Also, the generalized Tully-Fisher relationship for RFGC galaxies was finalized. It included data not only about HI linewidth and angular diameter in red and blue bands, but also about the morphological type of the galaxy and its surface brightness index. By that time the authors had information about radial velocities and HI $21\\,cm$ line widths or $V_{rot}$ for $1327$ RFGC galaxies from the different sources listed in the paper \\citep{ref:K00}. From this number, $1271$ galaxies were included in the sample, and the rest of them were considered to be outliers. As a result, the first list of peculiar velocities of RFGC galaxies was prepared by \\citet{ref:list}. Four years later, the number of galaxies with available data increased and reached $1561$ \\citep{ref:ParTug04}; $1493$ of them entered the sample. A new version of the list of peculiar velocities was prepared by \\citet{ref:ParTug05}. This list was the basis for solving the two abovementioned problems, namely mapping the matter density and estimation of cosmological parameters. The distribution of matter density was obtained in the paper \\citep{ref:SharPar} up to $75h^{-1}\\,Mpc$ in the supergalactical plane and 8 more planes. In the same article, the excess masses of attractors in this region were estimated, and the estimation of $\\beta$ parameter was obtained. The estimations of the cosmological parameters $\\Omega_m$ and $\\sigma_8$ were obtained in the paper \\citep{ref:cosm}. The model was expanded taking into account the effects of general theory of relativity in the paper \\citep{ref:ParGayd05}.  An important result, obtained in the paper \\\\ \\citep{ref:Par01}, and confirmed by \\\\ \\citet{ref:ParTug04}, was the statistical significance of quadrupole and octopole components (both at more than $99$ per cent confidence level). This result relied on the F-test, which assumes normal distribution of errors \\citep[see][]{ref:Hudson}. However, even the normally distributed errors of angular diameters and HI linewidths and deviations from Tully-Fisher relationship lead to non-Gaussian distribution of peculiar velocities. Thus, the statistical significance of these terms required additional consideration. In the article \\citep{ref:ParPar08} it was shown using Monte-Carlo simulations that the statistical significance of these components appeared to be less than perceived from the F-test, but still large enough to be considered: the quadrupole was statistically significant at $96$ per cent confidence level and the octopole -- at $90..93$ per cent confidence level.  In the four years that passed since the last list of peculiar velocities of RFGC galaxies was compiled, not only new data have appeared, but also some old data were remeasured. This led to a necessity of repeating all the steps required to obtain the new list. This means that not only new data have to be added, but the process of sample selection must be repeated from the very beginning. As it will be shown further, some results appeared to be notably different from the ones obtained with the previous sample.  In this paper we present the parameters of the large-scale collective velocity field in the framework of the multipole model as well as the estimations of some cosmological parameters. The distribution of matter density will be discussed in a separate article. ",
        "conclusions": "\\label{s:Conclusion}  We prepared a new increased and largely revised sample of RFGC galaxies with data about redshifts and HI linewidths. It allowed us to improve the estimations of parameters of the radial velocity field of large-scale collective motions of galaxies using 3 models of its multipole structure. In comparison with the previous versions the main features remained intact, but separate details, for example the apex of the dipole component of the bulk motion in the D-model, significantly changed. As before, the statistical significance according to F-test of both the quadrupole and the octopole components is well above $99.5$ per cent. The exact parameters are given in the article for subsamples with maximum distances $100h^{-1}\\,Mpc$ and $80h^{-1}\\,Mpc$.  The obtained velocity of the bulk motion is in agreement with the $\\Lambda$CDM model expectation of $\\sim 200\\,km\\,s^{-1}$. The value of the bulk motion velocity attracted additional attention after the recent paper of \\citet{ref:WFH09} who obtained this value as $407 \\pm 81\\,km\\,s^{-1}$ at the same scale $100h^{-1}\\,Mpc$ as in present article. In this connection some authors immediately started speculating about the challenge to the $\\Lambda$CDM model and the necessity for inclusion of non-gravitational forces \\citep{ref:AWW09}.  For $1623$ galaxies we compiled a list of peculiar velocities which will be published in the nearest future. We also intend to use it to determine the distribution of the matter density (including dark matter) at the scales $75h^{-1}\\,Mpc$ using POTENT method. This list was also used to estimate the cosmological parameters $\\Omega_m$ and $\\sigma_8$. The obtained constraint is given in graphical form on Fig. \\ref{fig:6}. The best numerical constraint is given by a combination $S_8=(\\Omega_m/0.3)^{0.35} \\sigma_8=0.91\\pm 0.05$."
    },
    "0910/0910.3060_arXiv.txt": {
        "abstract": "{Solar-like oscillations have now been observed in several stars, thanks to ground-based spectroscopic observations and space-borne photometry. CoRoT, which has been in orbit since December 2006, has observed the star HD49933 twice. The oscillation spectrum of this star has proven difficult to interpret. } {Thanks to a new timeseries provided by CoRoT, we aim to provide a robust description of the oscillations in HD49933, i.e., to identify the degrees of the observed modes, and to measure mode frequencies, widths, amplitudes and the average rotational splitting.} {Several methods were used to model the Fourier spectrum: Maximum Likelihood Estimators and Bayesian analysis using Markov Chain Monte-Carlo techniques.} {The different methods yield consistent result, and allow us to make a robust identification of the modes and to extract precise mode parameters.  Only the rotational splitting remains difficult to estimate precisely, but is clearly relatively large (several $\\mu$Hz in size).} {} ",
        "introduction": "\\label{sec:intro}  Stars are now the objects of seismic studies after decades of similar studies for the Sun, thanks to the advent of space-borne photometric observations (e.g., MOST, CoRoT and Kepler) and extremely precise ground-based spectroscopic observations \\citep[for a complete review, see for e.g.][]{Aerts2008SoPH}. This applies in particular to stars presenting solar-like p modes (acoustic oscillations stochastically excited by convection), with CoRoT observations of such stars showing clearly individual peaks in the Fourier spectra \\citep[e.g.][]{michel08}. Among these stars, HD49933 has already been the target of asteroseismic campaigns. It was first observed spectroscopically from the ground for 10 nights \\citep{mosser05}. In 2007, a first photometric time series of 60 days was collected by CoRoT, followed by a new long run of 137 days in 2008. HD49933 is a F5 main sequence star with an apparent visual magnitude $m_V = 5.77$. It is hotter than the Sun \\citep[$T_{\\rm eff}$=6780\\,K or $T_{\\rm eff}$=6500\\,K,][]{bruntt08,bruntt09,ryabchi09} with an estimated mass of $\\sim 1.2M_{\\odot}$ \\citep{mosser05} and an estimated radius of $1.34 \\pm 0.06\\,R_{\\odot}$ \\citep{thevenin06}. The surface rotation ($v\\sin i$) was determined to be around 10\\,km\\,s$^{-1}$ \\citep{mosser05,solano05}. The surface rotation period has also been measured at $\\sim$\\,3.4 days, using the 60-day CoRoT timeseries \\citep{appourchaux08,deheuvels2008}, from the signatures of photospheric transiting active regions (e.g., spots) which give rise to a clear peak in the very low-frequency part of the Fourier spectrum.  The seismic interpretation of HD49933 has proven to be very difficult. \\citet{mosser05} could not isolate individual p modes in the Fourier spectrum of observed line-of-sight velocities but were able to find a regular pattern in the spectrum, which is the signature of the large frequency separation between modes of same degree $l$ (but increasing radial order $n$). The first CoRoT time-series was analysed by \\citet{appourchaux08}, and these data clearly show individual p-mode peaks in the Fourier spectrum. However, the peaks show large widths, making the interpretation less than straightforward: a given peak could be interpreted as being a closely spaced pair of $l$=0 and $l$=2 modes, or a single (but rotationally split) $l$=1 mode. Based on the modeling of the spectrum using a Maximum Likelihood Estimator fitting method, \\citet{appourchaux08} chose one of these two possible interpretations (hereafter called model A) based on the highest {\\em likelihood} of each model. This first time-series was the object of other studies. \\cite{appourchaux2009a} put into perspective the results of \\citet{appourchaux08}, showing that the likelihood ratio test does not give the probability of the hypothesis given the data, but only the significance of the data given the hypothesis. \\cite{benomar09}, who applied a Bayesian analysis to the same time series, could not definitely favour one interpretation (model A) over the alternate (model B), based on the {\\em whole} probability distribution of each model. \\citet{gruberbauer2009a}, using a Bayesian approach too, also consider the identification ambiguous. \\citet{gaulme2009} used a simpler Bayesian approach (Maximum A Posteriori, or MAP, approach). The most probable model they found corresponds to the same identification as \\citet{appourchaux08}. More recently, \\citet{mosser09b} proposed an empirical method to determine the identification of the modes. Its application considers the model B as the more likely when using the two datasets used here.  It should be noted that the case of HD49933 is quite different from the solar case: solar modes are very narrow in comparison. Their widths (1\\,$\\mu$Hz for the modes with the highest amplitudes) are much smaller than the small-frequency separations between $l$=0 modes of order $n$ and the neighbouring $l$=2 modes of order $n$-1 (being typically around 10\\,$\\mu$Hz for the Sun). Moreover, the star inclination angle, if small, tends to attenuate the visibility of mode components with azimutal order $m\\ne0$, making the mode identification even more difficult.  Here, we use the new long-run CoRoT observations of HD49933, together with the original shorter run timeseries (Fig.\\,\\ref{fig:LC}), in order to properly describe the acoustic oscillations of the star clearly visible in the Fourier spectrum (Fig.\\,\\ref{fig:spec}, Fig.\\,\\ref{fig:diag_ech0} and Fig.\\,\\ref{fig:zoomspec}).  \\begin{figure*} \\begin{center} \\includegraphics[angle=90,width=18.4cm,height=6.8cm]{13111A.ps} \\end{center} \\caption{Mean power spectrum using 3 time series of 60 days (solid grey line), and the fitted model (dashed line).} \\label{fig:spec} \\end{figure*} ",
        "conclusions": "\\label{sec:conc}  The two CoRoT observation runs on the star HD49933 -- the first 60-day run, and the more recent longer run -- have now provided enough data to resolve the identification of modes in the oscillation spectrum at a very high confidence level. This identification relied on several independent analyses.  The new data have also allowed us to improve the precision in the mode parameters, with fractional improvements being in the range from 40\\,\\% to 70\\,\\% depending on the parameter. It is now possible to determine precise mode frequencies for $l$=0 and 1 modes. The $l$=2 mode parameters are more difficult to estimate because of the overlap with the stronger, neighbouring $l$=0 modes. The widths and amplitudes of the modes are well determined, as is the inclination angle of the star. The rotational frequency splitting remains the only mode parameter that is poorly constrained. However, it is clearly much higher than the solar value, and similar or larger in size to the inverse of the surface rotation period. In summary, this new information on the acoustic oscillations of HD49933 now opens the possibility for detailed seismic modeling of the star."
    },
    "0910/0910.0580_arXiv.txt": {
        "abstract": "Double-peaked \\oiii\\, profiles in active galactic nuclei (AGNs) may provide evidence for the existence of dual AGNs, but a good diagnostic for selecting them is currently lacking. Starting from $\\sim$ 7000 active galaxies in SDSS DR7, we assemble a sample of 87 type 2 AGNs with double-peaked \\oiii\\, profiles. The nuclear obscuration in the type 2 AGNs allows us to determine redshifts of host galaxies through stellar absorption lines. We typically find that one peak is redshifted and another is blueshifted relative to the host galaxy. We find a strong correlation between the ratios of the shifts and the double peak fluxes. The correlation can be naturally explained by the Keplerian relation predicted by models of co-rotating dual AGNs. The current sample statistically favors that most of the \\oiii\\ double-peaked sources are dual AGNs and disfavors other explanations, such as rotating disk and outflows. These dual AGNs have a separation distance at $\\sim 1$ kpc scale, showing an intermediate phase of merging systems. The appearance of dual AGNs is about $\\sim 10^{-2}$, impacting on the current observational deficit of binary supermassive black holes with a probability of $\\sim 10^{-4}$ (Boroson \\& Lauer). ",
        "introduction": "Supermassive black hole binaries and dual AGNs are the natural result of the current $\\Lambda$CDM cosmological paradigm, in which galaxies grow hierarchically through minor or major mergers. However, evidence for such systems is elusive and confirmed cases are rare despite circumstantial evidence. Examples can be found in Kochanek et al. (1999), Junkkarinen et al. (2001), Ballo et al. (2004), Maness et al. (2004), Guainazzi et al. (2005), Rodriguez et al. (2006), Hudson et al. (2006), Barth et al. (2008). The best-known examples of dual AGNs are found in NGC 6240 (Komossa et al. 2003), EGSD2 J142033+525917 (Gerke et al. 2007), EGSD2 J141550+520929 (Comerford et al. 2009a), the $z=0.36$ galaxy COSMOS J100043+020637 (Comerford et al. 2009b). Ideally, analysis of a large statistical sample should be undertaken.  The narrow emission lines in AGNs are generally thought to be produced by clouds in the narrow line regions at a scale of $\\sim 1$ kpc, where they are photoionized by the central energy source (e.g. Netzer et al. 2006). It is common that the \\oiii$\\lambda$5007 line has an asymmetric blue wing (Boroson 2005, Komossa et al. 2008), which could be caused by winds and outflows. On the other hand, about 1\\% AGNs have double-peaked \\oiii\\, profiles. The origin of the double-peaked emission lines remains unclear, but they have been interpreted as evidence for bi-polar outflows in the early 1980s (Heckman et al. 1981, 1984; Whittle 1985a,b,c) or disk-shaped narrow line regions (Greene et al. 2005). Evidence based on objects mentioned above suggests dual AGNs as an alternative explanation of the double-peaked profiles. However, more robust observational evidence is needed.  In this Letter, we select type 2 AGNs from SDSS DR7. The obscuration of the nucleus in these sources allows us to determine the redshifts of the host galaxies through the stellar absorption lines and investigate the properties of the double peaks. We find a strong correlation between the ratios of the shifts and the double peak fluxes. This is most naturally explained by the Keplerian relation of co-rotating dual AGNs. ",
        "conclusions": "We find a striking correlation between the ratios of \\oiii shifts and the double peak fluxes in a sample of 87 AGNs from the SDSS. This correlation strongly favors the explanation of dual AGNs in light of the Kepler diagram and does not favor a disk and outflow origin. The results also show that dual AGNs at 1 kpc scale are common, strongly impacting on the observational deficit of binary SMBHs with a probability of $\\sim 10^{-4}$ as shown by Boroson \\& Lauer (2009). The present correlation in the Kepler diagram and the dynamical criterion offer an efficient way to select targets (e.g. full version of Table 1) for future image detections in optical band and X-rays.  Finally, redshifts of host galaxies are a vital parameter in determination of the position of sources in the Kepler diagram. Higher resolution spectra are extremely important to diagnose dual AGNs among the weak double-peaked sources. The method presented in this paper is ideal for the study of such sources using high resolution spectra in the future."
    },
    "0910/0910.2911_arXiv.txt": {
        "abstract": "We present our ground-based CCD observations of the close binary systems DD~Mon and XY~UMa in B, V, R and I bands. The light curves are analyzed using the Wilson-Devinney code (W-D) for the derivation of the geometric and photometric elements of the systems. We compare the methods of photometric and spectroscopic mass ratio determination in these binaries, as a function of all typical difficulties, which arise during the analysis of such systems (light curve asymmetries, third light etc). Finally, a new spot model is suggested for the eclipsing system XY~UMa, which belongs to the RS~CVn type of active binaries. ",
        "introduction": "The B, V, R and I-band observations of both systems were carried out by means of CCD differential photometry. DD~Mon was observed on March 10-13, 2005 with the 1.22m Cassegrain reflector at the Kryoneri Astronomical Station of the National Observatory of Athens, Greece, while XY~UMa was observed on December 1, 4, 5 and 8, 2006 with the 0.40m Cassegrain reflector at the University of Athens Observatory, Greece. The light curves were analyzed by using the PHOEBE 0.29d software \\citep{PZ05}, which utilizes the W-D code \\citep{WD71,W79}.  For the photometric light curve analysis, we performed a q-search on both systems in modes 2, 4, 5, for a rough estimation of the photometric mass ratio ($q_{ph}$). We chose the range of ${0<q<1}$, and the best value (minimum sum of the square residuals - $\\chi^{2}$) of $q_{ph}$ was later used for the final solution (see Tables 2,~3, Fig.1). In both cases, the photometric mass ratios ($q_{ph}=0.575$ for DD~Mon and $q_{ph}=0.484$ for XY~UMa) were found to be smaller than the one obtained spectroscopically, which are $q_{sp} = 0.670(19)$ for DD~Mon \\citep{P09} and $q_{sp} = 0.70$ for XY~UMa \\citep{P07}.  The geometric and physical elements for both systems, together with the above radial velocities, are used to compute the absolute physical parameters (radii, masses and luminosity) in solar units (Table 1). In this table, we also include the solutions, derived with the use of the spectroscopically determined mass ratio, utilizing the (directly observed) radial velocities K1 and K2 as fixed parameters.    \\begin{table*} \\caption{Comparison of the absolute elements of DD~Mon and XY~UMa, obtained with photometrically and spectroscopically calculated mass ratio.} \\begin{flushleft} \\begin{small} \\scalebox{0.95}{ \\begin{tabular}{lcccccc}  \\hline \\hline \\textbf{Elem.}&    \\multicolumn{3}{c}{\\textbf{DD Mon}}   &   \\multicolumn{3}{c}{\\textbf{XY UMa}}    \\\\ &$q_{ph}=0.575$&$q_{sp}=0.670$&$diff.~(\\%)$&$q_{ph}=0.484$&$q_{sp}=0.610$&$diff.~(\\%$)\\\\ \\hline $M_{1}$       &    1.94(9)   &    1.39(7)   &    32.9    &    1.85(3)   &    1.13(8)   &    47.8    \\\\ $M_{2}$       &    1.12 (6)  &    0.93(6)   &    17.8    &    0.89(2)   &    0.69(4)   &    26.1    \\\\ $R_{1}$       &    1.62(3)   &    1.44(3)   &    12.1    &    1.37(1)   &    1.14(3)   &    17.8    \\\\ $R_{2}$       &    1.37(3)   &    1.29(3)   &    5.6     &    0.73(1)   &    0.68(1)   &    6.8     \\\\ $L_{1}$       &    3.54(13)  &    2.78(11)  &    24.1    &    1.32(2)   &    0.92(5)   &    35.3    \\\\ $L_{2}$       &    1.21(5)   &    1.08(5)   &    11.6    &    0.098(3)  &    0.093(4)  &    5.2     \\\\ $M_{bol1}$    &    3.36(4)   &    3.63(4)   &    7.5     &    4.44(1)   &    4.83(5)   &    8.4     \\\\ $M_{bol2}$    &    4.53(5)   &    4.66(5)   &    2.7     &    7.25(3)   &    7.31(5)   &    0.8     \\\\ \\hline \\hline  \\end{tabular}} \\end{small} \\end{flushleft} \\begin{small} \\end{small} \\end{table*}   \\begin{table*}  \\caption{The parameters of DD Mon derived from the LCs solution.} \\scalebox{0.6}{ \\begin{tabular}{lcccc|cccc} \\hline \\hline \\textbf{Parameter}          &       \\multicolumn{4}{c}{mode 5, adjusted q}      &            \\multicolumn{4}{c}{mode 5, fixed q}     \\\\ \\hline $\\varphi_{0}$               &            \\multicolumn{4}{c}{-0.0004(1)}         &                \\multicolumn{4}{c}{-0.0004(1)}      \\\\ $q~(m_{2}/m_{1})$           &            \\multicolumn{4}{c}{0.575(3)}           &           \\multicolumn{4}{c}{0.670}                \\\\ $i$ [deg]                   &            \\multicolumn{4}{c}{78.2}               &               \\multicolumn{4}{c}{77.2(1)}          \\\\ $T_1$$^{*}$[K], $T_2$ [K]   &          \\multicolumn{4}{c}{6250, 5202(4)}        &           \\multicolumn{4}{c}{6250, 5195(4)}        \\\\ $A_1$$^*$=$A_2$$^*$         &            \\multicolumn{4}{c}{0.5}                &                   \\multicolumn{4}{c}{0.5}          \\\\ $g_1$$^*$=$g_2$$^*$         &           \\multicolumn{4}{c}{0.32}                &               \\multicolumn{4}{c}{0.32}             \\\\ $\\Omega_{1}$, $\\Omega_{2}$  &        \\multicolumn{4}{c}{3.230(10), 3.037}       &       \\multicolumn{4}{c}{3.410(4), 3.190}          \\\\ &     $B$    &     $V$    &     $R$    &     $I$    &     $B$    &     $V$    &     $R$    &     $I$     \\\\ $L_{1}/L_{T}$               & 0.757(2)   &  0.711(2)  &  0.685(3)  &   0.658(3) &   0.730(2) &   0.685(2) &    0.659(2)&  0.633(3)   \\\\ $L_{2}/L_{T}$               &0.1959(3)   &  0.2242(4) & 0.2438(7)  &  0.2636(7) &   0.2133(3)&  0.2443(4) &   0.2657(5)&  0.2874(8)  \\\\ $L_{3}/L_{T}$               &0.047(2)    & 0.065(2)   &   0.072(3) &   0.079(4) &   0.056(2) &  0.071(2)  &   0.075(3) &  0.080(4)   \\\\ $X_{1}$, $X_{2}$            &0.677, 0.832&0.547, 0.689&0.468, 0.595&0.390, 0.501&0.677, 0.832&0.547, 0.690&0.468, 0.595&0.390, 0.502 \\\\ \\hline &   $pole$   &  $point$   &   $side$   &   $back$   &   $pole$   &  $point$   &   $side$   &   $back$    \\\\ \\hline $r_{1}$                     &  0.3702(1) & 0.4292(24) & 0.3867(12) & 0.4055(15) & 0.3598(5)  & 0.4178(12) & 0.3752(6)  &  0.39448)   \\\\ $r_{2}$                     &  0.3123(5) & 0.4452(8)  & 0.3263(5)  & 0.3586(5)  & 0.3231(6)  & 0.4589(10) & 0.3380(6)  &  0.3700(8)  \\\\ \\hline $\\chi^{2}$                  &             \\multicolumn{4}{c}{0.4221}            &                \\multicolumn{4}{c}{0.4327}          \\\\ \\hline\\hline & \\multicolumn{4}{c}{\\textit{$^*$assumed}, \\textit{$L_{T}=L_{1}+L_{2}+L_{3}$}}&\\multicolumn{4}{c}{}      \\\\ \\end{tabular}} \\end{table*}    \\begin{table*}  \\caption{The parameters of XY UMa derived from the LCs solution.} \\scalebox{0.57}{ \\begin{tabular}{lcccc|cccc} \\hline \\hline \\textbf{Parameter}          &          \\multicolumn{4}{c}{mode 5, adjusted q}   &       \\multicolumn{4}{c}{mode 5, fixed q}          \\\\ \\hline $\\varphi_{0}$               &            \\multicolumn{4}{c}{-0.0024(1)}         &           \\multicolumn{4}{c}{-0.0024(1)}           \\\\ $q~(m_{2}/m_{1})$           &            \\multicolumn{4}{c}{0.484(2)}           &               \\multicolumn{4}{c}{0.610}            \\\\ $i$ [deg]                   &            \\multicolumn{4}{c}{80.6(1)}            &               \\multicolumn{4}{c}{77.6(1)}          \\\\ $T_1$$^{*}$[K], $T_2$ [K]   &          \\multicolumn{4}{c}{5310, 3889(6)}        &           \\multicolumn{4}{c}{5310, 3806(6)}        \\\\ $A_1$$^*$=$A_2$$^*$         &            \\multicolumn{4}{c}{0.5}                &            \\multicolumn{4}{c}{0.5}                 \\\\ $g_1$$^*$=$g_2$$^*$         &           \\multicolumn{4}{c}{0.32}                &           \\multicolumn{4}{c}{0.32}                 \\\\ $\\Omega_{1}$, $\\Omega_{2}$  &        \\multicolumn{4}{c}{3.190(5), 3.670(10)}    &       \\multicolumn{4}{c}{3.425(3), 3.983(10)}      \\\\ &     $B$    &     $V$    &     $R$    &    $I$     &     $B$    &     $V$    &     $R$    &   $I$       \\\\ $L_{1}/L_{T}$               &   0.939(3) &  0.940(2)  & 0.928(2)   &  0.865(2)  &   0.947(3) &   0.940(3) &   0.924(2) &  0.869(2)   \\\\ $L_{2}/L_{T}$               &   0.0317(1)&  0.0475(1) & 0.0607(1)  &  0.0721(1) &   0.0343(1)&  0.0523(1) &  0.0678(1) &  0.0833(1)  \\\\ $L_{3}/L_{T}$               &   0.030(3) &  0.013(2)  &   0.012(1) &   0.063(1) &   0.019(2) &  0.008(2)  &  0.008(2)  &  0.047(1)   \\\\ $X_{1}$, $X_{2}$            &0.849, 0.823&0.787, 0.796&0.720, 0.770&0.633, 0.683&0.849, 0.828&0.787, 0.800&0.720, 0.770&0.633, 0.682 \\\\ \\hline &   $pole$   &  $point$   &   $side$   &   $back$   &   $pole$   &  $point$   &   $side$   &   $back$    \\\\ \\hline $r_{1}$                     & 0.3650(7)  & 0.4082(12) & 0.3795(8)  & 0.3935(9)  & 0.3503(4)  & 0.3937(6)  & 0.3635(4)  & 0.3785(5)   \\\\ $r_{2}$                     & 0.1994(10) & 0.2086(13) & 0.2019(11) & 0.2068(12) & 0.2135(7)  & 0.2234(9)  & 0.2163(8)  & 0.2214(8)   \\\\ \\hline \\textbf{Spot Parameter}     & $star~No~1$&  $spot~1$  &  $spot~2$  &            & $star~No~1$&  $spot~1$  & $spot~2$   &             \\\\ \\hline $co-latitude$ [deg]         &            &  82.3(4)   &  81.6(2)   &            &            &  82.4(3)   &  82.2(3)   &             \\\\ $longitude$ [deg]           &            &  192.6(3)  &  139.2(1)  &            &            &  190.6(3)  & 133.5(1)   &             \\\\ $radius$ [deg]              &            &  10.9(1)   &  11.4(2)   &            &            &  10.9(1)   &  11.4(2)   &             \\\\ $temp.~factor$              &            &  0.759(4)  & 0.765(4)   &            &            & 0.722(3)   &  0.759(4)  &             \\\\ \\hline $\\chi^{2}$                  &              \\multicolumn{4}{c}{0.0890}           &              \\multicolumn{4}{c}{0.1559}            \\\\ \\hline\\hline & \\multicolumn{4}{c}{\\textit{$^*$assumed}, \\textit{$L_{T}=L_{1}+L_{2}+L_{3}$}}&\\multicolumn{4}{c}{}      \\\\ \\end{tabular}} \\end{table*}   \\begin{figure}[] \\plottwo{LC_DD_Mon.eps}{LC_XY_UMa.eps}  \\caption{The observed and synthetic light curves of DD Mon and XY UMa.} \\end{figure} ",
        "conclusions": "The determination of the orbit of the secondary component is a major issue on near contact and detached systems, where the temperature and luminosity difference between the two components is large. Such a difficulty affects the mass ratio determination, where the proximity and eclipse effects play an additional role \\citep{ND03,VHW85}.  The radial velocities measured on XY~UMa components are affected by the presence of this effect. Therefore, the hotter area of the secondary is observed in smaller wavelengths and therefore smaller radial velocity is measured. This effect is also noticed by \\citet{P09}, who observed XY~UMa on the Mg triplet region of the spectrum and found a mass ratio of $q_{sp} = 0.70$, which is significantly larger than the the value of $q_{sp} = 0.61$ \\citep{PJ98}, found by infrared spectroscopy.  In our study we found that the photometric mass ratios for both systems were significantly smaller than the spectroscopic ones, which might be a result of the reflection effect. The geometric and physical elements describe very well the systems, using either the photometric or the spectroscopic mass ratio and the two theoretical models are distinguished only by the residual levels. However, the calculated absolute elements resulted in larger masses, radii and luminosities for both components. We conclude that the solutions based only on photometric data may hide such effects, giving solutions which might be even by 50\\% different than those expected."
    },
    "0910/0910.0313_arXiv.txt": {
        "abstract": "{ In this small review we present the actual state the knowledge about weighting black holes. Black holes can be found in stellar binary systems in our Galaxy and in other nearby galaxies, in globular clusters, which we can see in our and nearby galaxies, and in centres of all well-developed galaxies. Range of values of their masses is wide and cover about ten orders of magnitude (not taking into account the hypothetic primordial black holes). Establishing the presence of black holes, and in particular the measurement of their mass is one on the key issues for many branches of astronomy, from stellar evolution to cosmology. ",
        "introduction": "Astrophysical black holes (BH) are customarily divided into three classes: stellar mass black holes, with masses of a few up to twenty-thirty solar masses, intermediate mass black holes (IMBH) with masses of a few hundreds to a few thousands of solar masses, and massive (or super-massive) black holes residing in galactic dynamical centres, with masses from $10^5$ solar masses up.  Determination of values of black hole masses is important for the three main reasons: (i) to distinguish small mass black holes from neutron stars (ii) to analyse the luminosity states of accreting black holes through the luminosity to the Eddington luminosity ratio (ii) to study the growth of central black holes in galaxies. All three aspects are important, the last one is a key issue in cosmology, in studies of the galaxy evolution and AGN feedback.  The methods used for mass determination can be broadly divided into the following three classes: \\begin{itemize} \\item {dynamical methods} \\item {spectra fitting methods} \\item {scaling methods} \\end{itemize} although such a division is not necessarily unique. For example, the reverberation method is rather an example of a dynamical method while the Broad Line Region (BLR) radius vs. luminosity relation combined with a line width fitting is somewhere in between the direct dynamical method and a simple scaling law. In this review we will include it as a scaling method.  There are several recent reviews covering the topic of black hole mass determination \\citep[e.g.][]{Casares07,Ziolko08,Wandel08,VestergaardRev2009}. Here we will not go into the details of the various methods and its caveats but instead we will provide the collection of laws with short description of applicability. The possibilities are broad so black hole mass estimate is possible {\\it nearly} for every object. Different methods use observational data from various energy bands from radio through NIR and optical to X-ray domain. Therefore, the black hole mass measurement is one of the nice examples of the multifrequency approach. ",
        "conclusions": "\\subsection{Galactic black holes}  The typical accuracy in the determination of the black hole masses is about a factor 2 or better. The best known black hole mass (microquasar GRO J1655-40, $M_{BH} = 6.3 \\pm 0.5$, \\citealt{Greene2001}) was achieved due to the dynamical method (accurate period determination, inclination from modelling of the lightcurve due to the companion ellipticity, mass ratio also from modelling the ellipticity). Overall, there are over 40 (24 confirmed BH) determinations of the black hole mass in such systems in our Galaxy and a few in the binaries in nearby galaxies like in LMC and M33 (e.g. eclipsing binary M33 X-7, with black hole mass of $15.65 \\pm 1.45 M_{\\odot}$; \\citealt{Orosz07}).  \\subsection{Intermediate mass black holes}  This is the hottest subject in the Art of black hole mass measurement. First, the exact range is not specified, so it is not clear whether relatively small but otherwise typical AGN (i.e. located at the centres of their host galaxies) black holes from the range $\\sim 10^5 - 10^6 M_{\\odot}$ (see sample of \\citealt{Desroches09}) belong to this class. Their existence does not cause much doubt (e.g. \\citealt{HoRev}). The on-going stellar dispersion measurements in hundreds of near-by galaxies (see for example \\citealt{Ho2009}) may bring more of them although measuring black holes with masses below $10^6 M_{\\odot}$ is very difficult.  More questionable is the existence of the black holes with the masses of hundreds-thousands of the solar mass. Determination of $M_{\\rm BH}$ in several ULX sources, which are hosted by galaxies taken from  Messier and NGC catalogues, and the discussion of the stellar mass versus IMBH interpretation can be found in \\citet{Zamp09} (see also references therein).  \\subsection{Non-Active galaxies or Weakly Active galaxies}  The best example is of course the Milky Way galaxy, with the recent mass determination (see Sect.~\\ref{subsect:DynaMeth}). The mass in other nearby galaxies is known less accurately (for example, triple nucleus of Andromeda hosts a black hole mass in P3 of $(1.1 - 2.3) \\times 10^8 M_{\\odot}$, \\citealt{bender05}) but mass estimate (usually through stellar dispersion measurement) is available for hundreds of objects \\citep{HoRev}.  \\subsection{Active Galactic Nuclei}  The number of known AGN is of order of $10^5$ (in SDSS DR7) and rising due to massive surveys. The best results are obtained from water maser discussed in Sect.~\\ref{subsect:DynaMeth}. Generally, mass determination for tens/hundreds/thousands of objects were performed by several authors \\citep[e.g.][]{WooUrry02,VesterPeter06,Vestergaard08,Shen08,Fine08,kelly09,wu09}.  Numerous large surveys supplemented with the black hole mass and the Eddington ratio determinations \\citep[e.g.][]{Hickox2009} the available scaling laws from the list discussed in Sect.~\\ref{subsec:scaling}. The Eddington ratio is the frequently determine using another scaling to derive the bolometric luminosity, for example from the optical monochromatic flux. \\citet{Richards06} (see their Fig.~12) analysed spectra energy distributions of quasar type 1 and gave a formula: \\begin{equation} L_{\\rm bol} = {\\rm BC}_{\\lambda} \\times \\lambda L_{\\lambda} \\ , \\end{equation} where the bolometric corrections BC$_{\\lambda}$ = 3.81, 5.15, 9.26 are for measurements of luminosity, $\\lambda L_{\\lambda}$, made at $\\lambda$= 1350, 3000 and 5100 \\AA, respectively.  It is also possible to use the near-IR monochromatic luminosity \\begin{equation} \\label{eq:bolo_IR} L_{\\rm bol} = 6.4 L(\\rm NIR) \\end{equation} \\citep{cao05} or from the broad band X-ray luminosity \\citep{Hopkins2007} \\begin{equation} \\label{eq:bolo_X} L_{\\rm bol} = (10-20) L_{\\rm X}(0.5-8 {\\rm keV}). \\end{equation}  More complex, luminosity-dependent corrections can be found in \\citep{Hopkins2007}. \\citet{Hickox2009} used the coefficient in Eq.~(\\ref{eq:bolo_IR}) by a factor of two lower than derived by \\citep{Hopkins2007} in order to have consistent results for AGN with both IR and X-ray measurements available.  In heavily obscured Seyfert 2 galaxies or radio galaxies the bolometric luminosity can be estimated from the [$\\ion{O}{iii}$] line luminosity: \\begin{equation} L_{\\rm bol} = {\\rm C} L([\\ion{O}{iii}]), \\end{equation} where ${\\rm C} = 87, 142 $ and $454$ for logarithmic luminosity ranges of [$\\ion{O}{iii}$] line 38 - 40, 40 - 42 and 42 - 44, correspondingly \\citep{lamastra}."
    },
    "0910/0910.0724_arXiv.txt": {
        "abstract": "We present the first results from the new Survey of Extragalactic Nuclear Spectral Energies (SENSE) sample of ``blazars''. The sample has been chosen with minimal selection effects and is therefore ideal to probe the intrinsic properties of the blazar population. We report evidence for negative cosmological evolution in this radio selected sample and give an outline of future work related to the SENSE sample. ",
        "introduction": "Traditionally, blazar samples are selected using methods that can introduce severe selection effects. This makes it hard to distinguish between intrinsic sample properties and trends that have been induced by the selection process. Selection effects can therefore lead to erroneous conclusions being drawn about a sample that will hinder progress in understanding the physics of blazars. We have compiled a new redshift-limited sample of ``blazars'', designed to be free from many of the biases that blazar samples tend to incur. ",
        "conclusions": ""
    },
    "0910/0910.5930_arXiv.txt": {
        "abstract": "{Knowledge of the total population of symbiotic stars in the Galaxy is important for understanding basic aspects of stellar evolution in interacting binaries and the relevance of this class of objects in the formation of supernovae of type Ia.} {In a previous paper, we presented the selection criteria needed to search for symbiotic stars in IPHAS, the INT \\ha\\ survey of the Northern Galactic plane. IPHAS gives us the opportunity to make a systematic, complete search for symbiotic stars in a magnitude-limited volume.} {Follow-up spectroscopy at different telescopes worldwide of a sample of sixty two symbiotic star candidates is presented.} {Seven out of nineteen S-type candidates observed spectroscopically are confirmed to be genuine symbiotic stars.  The spectral type of their red giant components, as well as reddening and distance, were computed by modelling the spectra.  Only one new D-type symbiotic system, out of forty-three candidates observed, was found. This was as expected (see discussion in our paper on the selection criteria). The object shows evidence for a high density outflow expanding at a speed $\\ge$65~\\kms.  Most of the other candidates are lightly reddened classical T Tauri stars and more highly reddened young stellar objects that may be either more massive young stars of HAeBe type or classical Be stars. In addition, a few notable objects have been found, such as three new Wolf-Rayet stars and two relatively high-luminosity evolved massive stars. We also found a helium-rich source, possibly a dense ejecta hiding a WR star, which is surrounded by a large ionized nebula.} {These spectroscopic data allow us to refine the selection criteria for symbiotic stars in the IPHAS survey and, more generally, to better understand the behaviour of different \\ha\\ emitters in the IPHAS and 2MASS colour-colour diagrams.} ",
        "introduction": "Symbiotic stars, the interacting binaries with the longest periods, are recognised as key objects in the study of different aspects of stellar evolution. These include the formation of supernovae of type Ia (Munari \\& Renzini 1992, Hachisu et al. 1999), the powering mechanism of supersoft X-ray sources (cf. Jordan et al. 1996) and the formation of jets (Tomov 2003) and of bipolar (planetary) nebulae (\\cite{c03}).  The testing of these possible roles, and in particular their ability to form SNe Ia, depends on the total population of symbiotic stars within galaxies. This figure is, however, poorly known, even in our own Galaxy where less than 200 systems are known (\\cite{b00}), out of a predicted total population that ranges from 3$\\times$$10^3$ (\\cite{a84}) up to 4$\\times$$10^5$ (\\cite{mcm03}).  To improve the situation, a comprehensive search for symbiotic stars in the part of the Galactic plane visible from the northern hemisphere was started (Corradi et al. 2008, hereafter paper I).  The search begins with analysing the data from the INT Photometric \\ha\\ Survey of the northern Galactic plane (IPHAS, \\cite{d05}), taking advantage of the generally strong \\ha\\ emission and presence of a cool giant that characterises this class of objects.  Combination of these properties with constraints from 2MASS near--IR data defines the selection criteria adopted in paper I. There, a list of about one thousand symbiotic star candidates was drawn up from the set of IPHAS \\ha\\ emitters available at that time (Witham et al. 2008). As shown in paper I, this list suffers from significant contamination by reddened Be stars and young stellar objects, which are frequent in the Galactic plane.  In this paper, we present follow-up spectroscopy for a sample of 62 IPHAS candidate symbiotic stars. Observations are presented in Sect.~2. Section~3 is a discussion of the spectroscopically confirmed new symbiotic systems, while Sect.~4 outlines the majority group of mimics revealed in our spectra: among these there are several notable, rare object types. Discussion and perspectives are presented in Sect.~5.    \\begin{figure*}[!ht] \\centering \\includegraphics[width=17cm]{S182906.eps}\\\\[5pt] \\includegraphics[width=17cm]{S183501.eps}\\\\[5pt] \\includegraphics[width=17cm]{S184446.eps} \\caption{Spectra of the new S-type symbiotic stars IPHASJ182906.08--003457.2 (top), IPHASJ183501.83+014656.0 (middle), and IPHASJ184446.08+060703.5 (bottom). The blue and red regions of the spectra are displayed with different intensity cuts to show both faint and bright features.} \\label{F-specsymbio1} \\end{figure*}  \\begin{figure*}[!ht] \\centering \\includegraphics[width=17cm]{S184733.eps}\\\\[5pt] \\includegraphics[width=17cm]{S185323.eps}\\\\[5pt] \\includegraphics[width=17cm]{S193436.eps} \\caption{Spectra of the new S-type symbiotic stars IPHASJ184733.03+032554.3 (top), IPHASJ185323.58+084955.1 (middle), and IPHASJ193436.06+163128.9 (bottom).} \\label{F-specsymbio2} \\end{figure*}  \\begin{figure*}[!ht] \\centering \\includegraphics[width=17cm]{S193501.eps}\\\\[5pt] \\includegraphics[width=17cm]{S194607.eps} \\caption{Spectra of the S-type new symbiotic star IPHASJ193501.31+135427.5 (top), and of the D-type new symbiotic star IPHASJ194607.52+223112.3 (bottom).} \\label{F-specsymbio3} \\end{figure*} ",
        "conclusions": "\\label{S-diagram}  \\begin{figure*}[!ht] \\centering \\includegraphics[width=14cm]{figxx.eps} \\caption{IPHAS (top) and 2MASS (bottom) colour-colour diagrams for the objects observed spectroscopically. The three new IPHAS symbiotic stars from paper I are also included.  The locus of unreddened main-sequence and RGB stars are indicated by the solid and dotted lines, respectively. The arrow indicates the reddening vector for M0 giants: its length corresponds to 3~mag extinction in V.  The symbiotic stars selection boxes as defined in paper I are also indicated. In the IPHAS diagram, the boxes are common for both types. In the 2MASS diagram, the box for the D-types is to the right side of the contiguous box for the S-types.} \\label{F-cc} \\end{figure*}  The IPHAS and 2MASS colour-colour diagrams for the sample of sources observed spectroscopically, including the three new symbiotic stars presented in paper I, are shown in Fig.~\\ref{F-cc}.  The new symbiotic stars are generally displaced toward the upper-right side of the IPHAS diagram compared to the other \\ha\\ emitters that we have observed. In the 2MASS diagram, the S-type symbiotic stars form an isolated clump corresponding to reddened cool giants. T Tauri and other young stars are much more dispersed in both diagrams, and are generally found at smaller \\rha\\ colours than symbiotic systems.  Comparison with Figs.~1 and 2 of paper I shows that the new symbiotic stars from IPHAS generally have redder colours -- both in the optical and in the near-IR -- than the sample of objects of this class which were previously known. They are also optically fainter: their $r$ magnitudes are between 14.0 and 18.5, compared to the sample of known objects considered in paper I, which were mostly filling the magnitude interval between r=9.5 and r=14.5 mag. The luminosity difference is smaller if the K band is considered, which encourages the conclusion that the redder colours and fainter optical magnitudes are mainly due to the greater interstellar reddening, encountered along many of the lines of sight in the Galactic plane explored by IPHAS. Thus, as expected IPHAS seems to mainly discover reddened objects, which previously escaped detection because of their relatively faint optical magnitude. The IPHAS symbiotic stars discovered so far are concentrated between R.A.=18.5~h (close to the southern limit where the IPHAS coverage starts) and R.A.=20.5~h, i.e. at Galactic longitudes between 30 and 81 degrees. These are directions towards the inner Galaxy: considering the distances in Tab.~\\ref{T-symbio}, we conclude that at least half of the new symbiotic stars are located on the side of the bulge closer to the Sun.  Together with the three objects of paper I, so far we have found eleven new symbiotic systems in the Galactic plane. Considering that before our survey only eleven systems were known in the IPHAS area (Belczy\\'nski et al. 2000), we have doubled the number of symbiotic stars known in this region of the sky.  The result is positive, but we are still orders of magnitude below the large number of objects predicted to exist in the Galaxy (see Sect. 1). In order to obtain a new, reliable estimate of the total Galactic population of symbiotic stars, the following steps have to be followed:\\\\[-18pt] \\begin{itemize} \\item Build a complete sample of candidates.\\\\ The possible symbiotic stars in paper I were selected from the list of IPHAS \\ha\\ emitters by \\cite{w08}, which in turn was built when the survey's photometric catalogue was only partially completed.  We now need to produce a complete, magnitude limited sample of candidates, scanning the whole IPHAS photometric catalogue in a similar way as done by Viironen et al. (2009b) for compact PNe.  A similar search is foreseen in the opposite hemisphere when the IPHAS southern extension, VPHAS+, will be carried out at the 2.6mVST in Chile. The latter will also cover part of the Galactic bulge, where a large fraction of symbiotic stars are located. \\item Estimate accurately our success rate for the detection of symbiotic stars from the photometric selection.\\\\ So far, including the results in paper I, our success rate for S-type symbiotic stars is close to 50\\%\\ (10 out of 22 candidates observed spectroscopically). As paper I lists 337 candidates, we might expect to find another 150 new S-type systems if we continue the spectroscopic follow up. However, the newly discovered systems have redder \\ri\\ and \\jh\\ colours than the majority of the candidates (cf. Fig. 6 of paper I), most of which are also quite faint.  Thus if the remaining objects prove to be symbiotic stars, they would be faint red giants with little reddening, i.e. faraway objects in clear lines of sight, a combination that is unlikely in the Galactic plane. Indeed, all the objects located in the lowest part of the selection box for S-types observed spectroscopically, turned out to be young stars. We conclude that our global success rate will be significantly lower than obtained so far. More spectra are needed in order to determine this.  We have found only one new D-type system. Spectroscopy confirms the idea discussed in paper I that the vast majority of candidates in the adopted D-type selection boxes are young or massive stars. As D-type systems represent 15--20\\%\\ of the known sample of Galactic symbiotic stars, the effect of missing them is not dramatic and can be taken into account via a correction factor. \\item Transform the estimate of the total number of objects in the IPHAS area to the total number in the Galaxy.\\\\ This implies introducing the results of our search (completeness and success rate as a function of magnitude and direction) into a model that includes the spatial distribution of symbiotic stars in the Galaxy and the growth of interstellar extinction in every direction. \\end{itemize}"
    },
    "0910/0910.4985_arXiv.txt": {
        "abstract": "A quintessence scalar field or cosmon interacting with neutrinos can have important effects on cosmological structure formation. Within growing neutrino models the coupling becomes effective only in recent times, when neutrinos become nonrelativistic, stopping the evolution of the cosmon. This can explain why dark energy dominates the Universe only in a rather recent epoch by relating the present dark energy density to the small mass of neutrinos. Such models predict the presence of stable neutrino lumps at supercluster scales ($\\sim 200$ Mpc and bigger), caused by an attractive force between neutrinos which is stronger than gravity and mediated by the cosmon. We present a method to follow the initial nonlinear formation of neutrino lumps in physical space, by integrating numerically on a 3D grid nonlinear evolution equations, until virialization naturally occurs. As a first application, we show results for cosmologies with final large neutrino average mass $\\sim 2$ eV: in this case, neutrino lumps indeed form and mimic very large cold dark matter structures, with a typical gravitational potential $10^{-5}$ for a lump size $\\sim 10$ Mpc, and reaching larger values for lumps of about $200$ Mpc. A rough estimate of the cosmological gravitational potential at small $k$ in the nonlinear regime, $\\Phi_{\\nu} = 10^{-6} (k/k_0)^{-2}\\,,\\,\\, 1.2\\cdot10^{-2}$ h/Mpc $< k_0 < 7.8\\cdot10^{-2}$ h/Mpc, turns out to be many orders of magnitude smaller than an extrapolation of the linear evolution of density fluctuations. The size of the neutrino-induced gravitational potential could modify the spectrum of CMB anisotropies for small angular momenta. ",
        "introduction": "The presence of an interaction between a quintessence scalar field or cosmon and other species in the Universe \\cite{wetterich_1995} \\cite{amendola_2000} influences the nature and properties of dark energy, with relevant effects on structure formation \\cite{perrotta_baccigalupi_2002, nusser_etal_2005, koivisto_2005, kesden_kamionkowski_2006,mainini_bonometto_2006, farrar_rosen_2007, bean_etal_2007, cabral_etal_2009, keselman_etal_2009, afshordi_etal_2005, lavacca_etal_2009,bernardini_etal_2009}. Recently, {\\emph N}-body simulations have been performed for dark matter particles interacting with dark energy \\cite{baldi_etal_2008, maccio_etal_2004, li_zhao_2009}. We concentrate here on growing neutrino models \\cite{amendola_etal_2007, wetterich_2007}, where the neutrino-cosmon coupling can explain the ``why now?'' problem of dark energy. It has been recently shown \\cite{mota_etal_2008} that this predicts the existence of very large (supercluster) structures.  Indeed, the key ingredient of growing neutrino quintessence is the pre\\-sence of a coupling between dark energy and neutrinos. The latter have a mass that grows with the evolution of the Universe - $m_\\nu(\\phi)$ being a function of the cosmon field $\\phi$. The cosmon-neutrino coupling $\\beta$ is given by the logarithmic derivative $\\beta = -d\\ln{m_\\nu}/d \\phi$. Since neutrino masses become cosmologically relevant only for redshift $z \\approx 5$, this framework can naturally answer the question why only in recent times dark energy leads the Universe expansion to accelerate, thus providing a solution to the coincidence problem.  In these models, as long as neutrinos stay relativistic, the coupling plays no role and dark energy tracks the background, along the attractor trajectories characterizing the cosmon evolution in the presence of an exponential potential \\cite{wetterich_1995, wetterich_1988, ratra_peebles_1988, ferreira_joyce_1998, copeland_etal_1998}. When neutrinos become nonrelativistic, the coupling between neutrinos and quintessence becomes relevant and almost stops the evolution of the cosmon. Then dark energy starts to resemble a cosmological constant with roughly the value of the exponential potential at the end of the attractor era. The transition from the attractor solution to an almost static solution is therefore strictly connected to a cosmological event, that is neutrinos becoming nonrelativistic. This ``trigger event'' leads dark energy to dominate over cold dark matter, naturally starting the recent era of accelerated expansion.  Growing neutrino quintessence requires a cosmon-neutrino interaction somewhat stronger than gravity - typically the attraction between nonrelativistic neutrinos exceeds gravity by a factor $10^{3}$. This coupling is substantially larger than a possible cosmon coupling to atoms. This may be motivated by the particular particle physics mechanism responsible for the neutrino mass, which typically involves a heavy singlet field and not only the standard Higgs doublet \\cite{wetterich_2007}. Even much larger neutrino couplings leading to a strongly coupled ``acceleron-neutrino fluid'' have been investigated within mass varying neutrino models \\cite{fardon_etal_2004, afshordi_etal_2005, bjaelde_etal_2007}. In particular, mass varying neuitrino models employ a scalar field with a mass much larger than the Hubble parameter. For growing neutrino models, in contrast, the time dependent cosmon mass equals the Hubble parameter up to a factor of order one, similar to many models of coupled quintessence. (For coupled quintessence with neutrinos at the linear level see ref. \\cite{bean_etal_2007, mota_etal_2008, brookfield_etal_2005, ichiki_keum_2007, franca_etal_2009}.) For this reason, the cosmon and the neutrinos always have to be treated as separate ingredients rather than as a common fluid. We remark that the small time dependent cosmon mass follows naturally from possible explanations of an exponential potential in forms of an asymptotically vanishing dilatation anomaly \\cite{wetterich_2008}. It requires no additional fine tuning of particle physics parameters.  Furthermore, in growing neutrino cosmologies the coupling is ineffective for most of the cosmological evolution and only becomes active when neutrinos become nonrelativistic, relating naturally dark energy and neutrino properties. In view of bounds on the present neutrino mass, $m_\\nu(t_0) < 2.3 {\\rm eV}$ \\cite{amsler_etal_2008}, and the time dependence of $m_\\nu$, which makes the mass even smaller in the past, the time when neutrinos become nonrelativistic is typically in the recent history of the Universe, say $z_{{\\rm NR}} \\approx 5-10$ \\cite{amendola_etal_2007}. It is only from this time on that neutrinos start feeling the effects of the coupling, manifesting effectively as a new attractive interaction between neutrinos. The `fifth force' responsible for the formation of neutrino lumps ``switches on'' only in rather recent cosmology.  Neutrino fluctuations on length scales larger than the free streaming length are still present at $z_{{\\rm NR}}$, and they start growing for $z < z_{{\\rm NR}}$ with a large growth rate. As illustrated in \\cite{mota_etal_2008}, this opens the possibility that neutrino perturbations rapidly grow nonlinear on supercluster scales and beyond. Neutrino nonlinear fluctuations later turn into bound neutrino lumps of the type discussed in \\cite{brouzakis_etal_2007}, thus opening a window for observable effects of the growing neutrino scenario. The linear analysis \\cite{mota_etal_2008} has already provided an estimate of this effect as a function of redshift and scale. Typically large scale neutrino fluctuations with size $\\gsim 10$ Mpc become nonlinear at a redshift $z \\approx 1$.  As noted in \\cite{mota_etal_2008}, a continuation of the linear evolution beyond the time when the neutrino fluctuations are of order unity can easily produce erroneous results. In the linear approximation, the neutrino density contrast would quickly reach huge values, producing a very large gravitational potential. Then a very strong ISW effect would seemingly indicate a strong conflict with the cosmic microwave anisotropies (cf. ref. \\cite{franca_etal_2009} for a linear analysis). The true physics differs very strongly from the linear behavior, for example by the asymmetry between very large positive density contrasts, while a negative density contrast is bounded by $\\delta_{\\nu} \\geq -1$ since the neutrino density cannot be negative. In order to understand the true gravitational potentials that will be generated by the neutrino lumps, one has to understand the nonlinear dynamics of how local neutrino lumps form, how they are distributed in mass and size and how they may merge into larger lumps as time goes on.  In this work we investigate the formation of individual neutrino lumps at a nonlinear level and in the Newtonian limit. One may wonder if a spherical collapse approach suffices to give a meaningful description of the lumps. We find that this is not the case, as the major force driving neutrinos to collapse is not gravity but the additional fifth force introduced by the coupling to the cosmon and dominant once neutrinos decouple from the background expansion. The evolution of the effective cosmic scale factor for the space occupied by the neutrino lump is only a subleading effect, in contrast to the formation of dark matter halos. We have therefore developed a numerical method for solving the hydrodynamic equations for the neutrino fluid, coupled to the cosmon and gravity. We have also included dark matter, but this is a subleading effect.  The nonlinear analysis developed in this work provides a self-consistent way of analyzing the growth of neutrino perturbations in growing neutrino models. We show that neutrinos indeed form stable structures on large sub-horizon scales and we estimate the properties of the lump as a function of redshift and scale. In particular, we compute the gravitational potential of the lump at a time when the collapse ends due to virialization. This is a key quantity for an estimate of the effects of neutrino lumps on large scale structure - as large scale peculiar velocities - or on the cosmic microwave anisotropies in form of the integrated Sachs - Wolfe (ISW) effect.  This paper is organized as follows. In section \\ref{gnm} we recall the framework of growing neutrino cosmologies in which neutrino lumps form. In section \\ref{nl} we introduce the set of equations describing the evolution of neutrino overdensities at a nonlinear level. We comment on the methods used in the numerical integration and present our results for the case of large present neutrino masses in section \\ref{sec:results}. Section \\ref{sec:ini_conds} discusses the initial conditions used for the nonlinear analysis. In section \\ref{sec:rel_nonrel} we relate the nonlinear equations to relativistic linear ones. Finally we draw our conclusions in section \\ref{conclusions}. ",
        "conclusions": "\\label{conclusions} In growing neutrino cosmologies, the neutrino mass grows in time as a function of a light dark energy scalar field, the cosmon. The resulting coupling between cosmon and neutrinos modifies the evolution of the cosmon as soon as neutrinos become nonrelativistic. From this time on the cosmon evolves only very slowly, and its potential energy acts similar to a cosmological constant. This can provide a natural explanation of the coincidence problem without effectively introducing new parameters. The current dark energy density can be related to the neutrino mass. In addition, the interaction gives rise to a new long-range attractive force for neutrinos. This fifth force is responsible for a rapid growth of neutrino perturbations. These perturbations become nonlinear on large length scales and may eventually form very large stable structures, neutrino lumps.  We have extended previous studies on the topic to the nonlinear regime, providing a method to investigate the formation of neutrino lumps by evaluating Navier Stokes equations, in which the dragging term due to the coupling has been suitably included. As we have verified, these are the most general equations to compute the spacetime evolution of perturbations subject to external forces within the Newtonian limit. Unlike spherical collapse methods, the full nonlinear equations describe coherently the growth of structures due to an external force stronger than gravity. We note that although our calculations were performed within a growing neutrino scenario, the results are not limited to this case. Instead, the method can be applied for general coupled quintessence models, or whenever a force different from gravity is present, driving the growth of nonrelativistic matter perturbations.  We have numerically solved the set of hydrodynamical equations on a three-dimensional spatial grid. This has allowed us to follow the evolution of a single neutrino lump in physical space and to determine its properties as it approaches stability. Nonrelativistic neutrinos decouple from the background expansion and collapse into virialized stable structures. We have identified the characteristic gravitational potential of these structures as a function of redshift and lump size and we have illustrated the detailed behavior of the lumps from the linear regime to virialization. Indeed, after a period of fifth-force driven collapse, the build-up of large nonradial velocities leads to stabilization of the lumps. By solving on a three-dimensional grid with a general initial profile, we are able to provide a detailed illustration of the evolution at different distances from the center of the lump, following the change in kinetic and potential energy as well as the increasing contribution of nonradial velocities. We have also estimated the density profile of the virialized neutrino lumps and the associated profile of the gravitational potential. Limitations of our results for characteristic properties of simple neutrino lumps arise from two issues. First, at the present stage the numerical precision of our algorithm does not allow us to compute the behavior after virialization sets in and to follow the evolution until the neutrino lump becomes approximately stable. Second, our work concentrates on the evolution of single lumps while we have not addressed interesting topics on cosmological scales such as the dynamics of the lumps and their possible merging.  The cosmological gravitational potential provides an important mean to test the model versus observations and can be used to estimate the impact of neutrino lumps on the CMB angular power spectrum. An extrapolation from the gravitational potential for single lumps to the neutrino lump induced cosmological potential $\\Phi_{\\nu}(k)$ at a given wave number $k$ involves an averaging over the distribution of neutrino lumps. The present work can only give a very rough estimate of the result of this averaging. Nevertheless, it is apparent that the values of $\\Phi_{\\nu}(k)$ resulting from the nonlinear hydrodynamical equations remain many orders of magnitude below the extrapolation of the linearized equations if $k \\gg 10^{-3}$ Mpc$^{-1}$. We find a typical value $\\Phi_{\\nu}(k) = 10^{-6} (k/k_0)^{-2}$, where the large uncertainty is reflected by the uncertainty in $k_0$ that we estimate in the range $8.6\\cdot10^{-3}\\text{ Mpc}^{-1} < k_0 < 5.6\\cdot10^{-2}\\text{ Mpc}^{-1}$. This estimate has to be modified for $k \\lsim 4.5\\cdot10^{-3}$ Mpc$^{-1}$, where neglected effects beyond the Newtonian approximation become relevant, and for $k \\gsim 1.1$ Mpc$^{-1}$, where neutrino free streaming prevents the formation of lumps. Also, a formation history involving merging may reduce $\\Phi_{\\nu}(k)$, especially for small $k$ corresponding to very large lumps.  The characteristic size of $\\Phi_{\\nu}(k)$ found in this paper is of an order of magnitude where it may leave imprints on the CMB spectrum at small angular momenta. In particular, the late ISW effect could be enhanced, leading to stronger correlations between temperature fluctuations in the CMB and observed large scale structures in the gravitational potential. (The present observational situation hints to an enhancement of this correlation by a factor of about $2$ as compared to the $\\Lambda$CDM model \\cite{ho_etal_2008}.) Small oscillations of the CMB spectrum for angular momenta $l \\lsim 50$ are also conceivable. Without a more reliable estimate of $\\Phi_{\\nu}(k,z)$ it seems premature to judge if a given growing neutrino model remains compatible with observation. We recall in this context that the present work concentrates on a particular class of growing neutrino models (constant $\\beta$) and assumes a large average present neutrino mass $\\bar{m}_{\\nu}(t_0) \\sim 2$ eV. We expect that a smaller $m_\\nu$ reduces the cosmological effects of neutrino lumps since the total neutrino fraction of the energy density $\\Omega_\\nu$ gets reduced and the neutrino induced effects set in at even more recent cosmological times.  It may be possible to directly observe large neutrino lumps through their gravitational potential. A possible indication would be an observation of structures at large length scales where fluctuations in the $\\Lambda$CDM model are expected to remain linear such that substantial over- or underdensities are very rare. The gravitational potential of large neutrino lumps may influence a substantial fraction of space. If we are sitting not too far from a neutrino lump (or even within a neutrino lump) this may induce a certain anisotropy of the observed sky on the largest scales (i.e. small differences between northern and southern hemisphere or similar effects). Finally, the rapid recent growth of the neutrino induced gravitational potential could lead to an enhancement of the peculiar velocities as a characteristic effect of the nongravitational interactions \\cite{ayaita_etal_2009}. Present observations of the large scale bulk flow seem to suggest values substantially larger than expected in the $\\Lambda$CDM model \\cite{perivolaropoulos_2008,watkins_2008,kashlinsky_2008,lavaux_2008}. Finally, the cosmon may also have a coupling to dark matter, substantially smaller than the cosmon-neutrino coupling. In this case the fifth force effects could also influence the behavior of dark matter. Detection of such effects would challenge standard $\\Lambda$CDM models, providing a hint for interacting cosmologies such as growing neutrino models.  Let us end with the remark that there is a chance that the time variation of the neutrino mass could even be detected. Indeed, our understanding of structure formation and other cosmological features places strong upper bounds on the neutrino masses in early cosmology, say for $z \\gtrsim 5$. One may argue about the precise location of this bound, but an average neutrino mass $m_{\\nu}(z > 5)$ of $0.5$ eV would certainly have left a strong imprint on cosmology which has not been observed. In consequence, if the direct searches for a neutrino mass or the neutrinoless double beta decay indicate a present neutrino mass larger than $0.5$ eV, this could be interpreted as a strong signal in favor of a growing neutrino mass. \\\\ \\\\ \\\\"
    },
    "0910/0910.1044_arXiv.txt": {
        "abstract": "We present new tests for disruption mechanisms of star clusters based on the bivariate mass-age distribution $g(M, \\tau)$. In particular, we derive formulae for $g(M,\\tau)$ for two idealized models in which the rate of disruption depends on the masses of the clusters and one in which it does not. We then compare these models with our {\\it Hubble Space Telescope\\/} observations of star clusters in the Antennae galaxies over the mass-age domain in which we can readily distinguish clusters from individual stars: $\\tau \\la 10^7 (M/10^4 M_{\\odot})^{1.3}$~yr. We find that the models with mass-dependent disruption are poor fits to the data, even with complete freedom to adjust several parameters, while the model with mass-{\\it in\\/}dependent disruption is a good fit. The successful model has the simple form $g(M,\\tau) \\propto M^{-2} \\tau^{-1}$, with power-law mass and age distributions, $dN/dM \\propto M^{-2}$ and $dN/d\\tau \\propto \\tau^{-1}$. The predicted luminosity function is also a power law, $dN/dL \\propto L^{-2}$, in good agreement with our observations of the Antennae clusters. The similarity of the mass functions of star clusters and molecular clouds indicates that the efficiency of star formation in the clouds is roughly independent of their masses. The age distribution of the massive young clusters is plausibly explained by the following combination of disruption mechanisms: (1) removal of interstellar material by stellar feedback, $\\tau \\la 10^7$~yr; (2) continued stellar mass loss, $10^7 {\\rm yr} \\la \\tau \\la 10^8 {\\rm yr}$; (3), tidal disturbances by passing molecular clouds, $\\tau \\ga 10^8$~yr. None of these processes is expected to have a strong dependence on mass, consistent with our observations of the Antennae clusters. We speculate that this simple picture also applies---at least approximately---to the clusters in many other galaxies. ",
        "introduction": "Star clusters are born in molecular clouds and are then dispersed into the surrounding stellar field by a variety of processes operating on different timescales, including the removal of interstellar material by stellar feedback, continued stellar mass loss, tidal disturbances by passing molecular clouds, gravitational shocks during rapid passages near the galactic bulge or through the galactic disk, orbital decay into the galactic center caused by dynamical friction, and stellar escape driven by internal two-body relaxation. Signatures of these processes are encoded in the mass function, $\\psi(M) \\propto dN/dM$, and the age distribution, $\\chi(\\tau) \\propto dN/d\\tau$, for a population of clusters in a particular galaxy. A more informative, but less familiar statistic is the bivariate distribution of masses and ages, $g(M,\\tau) \\propto \\partial^2N/\\partial{M}\\partial{\\tau}$. The key feature of $g(M,\\tau)$ is that it includes correlations between $M$ and $\\tau$, information that is absent from the univariate distributions $\\psi(M)$ and $\\chi(\\tau)$. Thus, with the help of $g(M, \\tau)$, we can answer questions such as whether low-mass clusters are disrupted faster than, or at the same rate as, high-mass clusters. In this paper, we derive formulae for $g(M, \\tau)$ for the first time for several different models for the disruption of clusters, and we then compare these formulae with our {\\it Hubble Space Telescope\\/} ({\\it HST\\/}) observations of clusters in the Antennae galaxies. In a companion paper, we present similar comparisons for the clusters in the Large and Small Magellanic Clouds (LMC and SMC; Chandar et al. 2009, hereafter CFW09).  We have focused on the star clusters in the Antennae galaxies for several reasons. First, the population of clusters is large, $N \\approx 2300$ brighter than $M_V = -9$, ample for statistical studies. Second, many of the bright young clusters have masses in the range of old globular clusters, $10^4 \\la M/M_{\\odot} \\la 10^6$, suggesting that the former may simply be younger versions of the latter. Third, because the Antennae are currently in the throes of a major merger, they give us a close-up, internal view of events and processes that were more common in galaxies in the past, during their hierarchical formation. Fourth, the Antennae have been mapped at almost every wavelength available to astronomers, from radio to X-rays, providing a more detailed picture of their stellar and interstellar contents than for almost any other galaxy outside the Local Group.  In our previous studies of the clusters in the Antennae, we found that the mass and age distributions could be approximated by power laws: $\\psi(M) \\propto M^{\\beta}$ with $\\beta \\approx -2$ (Zhang \\& Fall 1999, hereafter ZF99) and $\\chi(\\tau) \\propto \\tau^{\\gamma}$ with $\\gamma \\approx -1$ (Fall et al. 2005, hereafter FCW05). In these studies, we also found that the shapes of $\\psi(M)$ and $\\chi(\\tau)$ are approximately independent of the age and mass limits, respectively, of the samples used to determine them. These findings suggest that the bivariate mass-age distribution can be approximated simply by the product of the univariate distributions: \\begin{equation} g(M, \\tau) \\propto \\psi(M)\\chi(\\tau) \\propto M^{\\beta}\\tau^{\\gamma} \\end{equation} \\begin{equation} {\\rm for} \\,\\,\\,\\,\\,\\,\\,\\,\\, \\tau \\la 10^7 (M/10^4 M_{\\odot})^{1.3}\\,{\\rm yr}. \\end{equation} The second expression above indicates the approximate domain of validity of the first, the condition that objects in the sample be brighter than the most luminous individual stars ($L_V \\ga 3\\times10^5 L_{\\odot}$).\\footnote {Equation~(2) follows from the fact that the mass-to-light ratios of clusters vary with age approximately as $M/L_V \\propto \\tau^{0.8}$ for $\\tau \\ga 10^7$~yr.} In this model, clusters form with a power-law initial mass function and are then disrupted at rates that are independent of their masses. We have already explored some of the consequences of equation~(1) in two recent papers (Fall 2006; Whitmore et al. 2007, hereafter WCF07). In this paper, we make further tests of this model, and we also compare our {\\it HST\\/} observations of the Antennae clusters with two models in which clusters of different masses are disrupted at different rates.  The age distribution for a population of star clusters reflects, in principle, a combination of both formation and disruption rates. We interpret the decline of $\\chi(\\tau)$ primarily as the result of disruption rather than formation for the following reasons. First, the age distribution has a sharp peak at the present time, $\\tau = 0$, and a small width, characterized by the median age $\\tau \\sim 10^7$~yr. This requires either rapid disruption at an arbitrary time (now) or rapid formation at a special time (now), the former being much more likely a priori than the latter. Of course, the Antennae galaxies are presently merging, which has almost certainly boosted the formation rates of stars and clusters, but this occurs more slowly, on the orbital timescale of the galaxies, $\\tau \\sim {\\rm few} \\times 10^8$~yr. In the simulations of merging galaxies by Mihos et al. (1993), including the Antennae, for example, the star formation rate varies by factors of only a few or less in the past $10^8$~yr, whereas the observed age distribution of the clusters drops by nearly 2 orders of magnitude in this interval of time. The second reason we interpret the decline of $\\chi(\\tau)$ in terms of disruption is that it has the same shape, including the sharp peak at $\\tau = 0$, in large regions separated by distances of order 10~kpc within the Antennae galaxies (Whitmore 2004; WCF07). There is no physical mechanism that could synchronize a burst of cluster formation this precisely at these separations; the communication time is $\\tau \\sim 10^9$~yr for a signal traveling at the typical random velocity in the interstellar medium (ISM; $v \\sim 10$~km s$^{-1}$) and $\\tau \\sim 10^8$~yr for one traveling at the highest bulk velocity ($v \\sim 100$~km s$^{-1}$). Thus, we conclude that the observed decline in the age distribution is caused mainly by disruption, at least for $\\tau \\la \\mbox{few}\\times10^8$~yr.  We have speculated that equation~(1) for $g(M, \\tau)$ may also describe---approximately---the cluster populations in other galaxies (Fall 2006; WCF07). If this picture turns out to be generally valid, it will mean that star clusters form and evolve in much the same way in different galaxies, the primary variable being an overall scale factor proportional to the total number of clusters and hence to their average formation rate. We define a cluster here to be any concentrated aggregate of stars with a density much higher than that of the surrounding stellar field, whether or not it is gravitationally bound, since this is virtually impossible to determine from the available observations, especially for clusters younger than about 10 internal crossing times. Most, if not all, stars form in clusters (as just defined) and most clusters then dissolve into the stellar field. We find it intriguing that this whole complex process might be represented approximately by a simple ``recipe'' such as equation~(1).  There are a few other hints that point toward this model: (1) The same mass and age distributions, $\\psi(M) \\propto M^{-2}$ and $\\chi(\\tau) \\propto \\tau^{-1}$, are found for clusters with $M \\la 10^3 M_{\\odot}$ and $\\tau \\la 3 \\times 10^8$~yr in the solar neighborhood (Lada \\& Lada 2003). (2) The luminosities of the brightest clusters in different galaxies scale approximately with the size of the population in the way expected from equation~(1) (Larsen 2002; Whitmore 2003; WCF07). (3) The mass spectrum of molecular clouds, from which the clusters form, appears to be similar in different galaxies (Blitz et al. 2007). (4) The most important early disruption processes, removal of ISM by stellar feedback and continued stellar mass loss, are independent of the properties of the host galaxy (except possibly its metallicity and stellar initial mass function (IMF)). Of course, a definitive conclusion about the generality of equation~(1) will require more detailed studies, such as that presented here for the Antennae, of the cluster populations in several other galaxies. We have recently made a start on this with the LMC and SMC (CFW09). Our main purpose in this paper is to present some of the interpretive tools needed for such studies.  The plan for the remainder of this paper is the following: In \\S2, we present analytic formulae for $g(M, \\tau)$ for two models with mass-dependent disruption and the model already mentioned with mass-independent disruption. In \\S3, we describe our {\\it HST\\/} observations of the Antennae clusters and compare these with the models presented in the previous section. We also derive the luminosity function of the Antennae clusters and show how this is related to the mass function through $g(M, \\tau)$. In \\S4, we discuss the physical processes most likely responsible for the formation and disruption of clusters and how these processes affect the mass and age distributions. This interpretative section presents our current understanding of the entire life cycles of star clusters. In \\S5, we summarize our main results and their broader implications. The paper also includes two appendices. The first (A) discusses some related work on the mass and luminosity functions of the Antennae clusters by Mengel et al. (2005) and by Anders et al. (2007) respectively, while the second (B) presents some specialized formulae needed in \\S3. ",
        "conclusions": "In this paper, we have derived analytical formulae for the bivariate mass-age distribution $g(M,\\tau)$ for three idealized models of the formation and disruption of star clusters. We have also derived formulae for the corresponding averages of $g(M, \\tau)$ over finite intervals of $M$ and $\\tau$, denoted by $\\bar{g}(\\tau)$ and $\\bar{g}(M)$ respectively. In all three models considered here, clusters are assumed to form with a power-law initial mass function at a constant rate. In the first two models, proposed by Boutloukos \\& Lamers (2003), clusters are disrupted on a timescale that depends on their masses as $\\tau_d(M) \\propto M^k$ with $k > 0$, either suddenly (Model~1) or gradually (Model~2). In the third model, clusters are disrupted (suddenly or gradually) at a fractional rate independent of their masses, as indicated by our earlier studies of the Antennae clusters (ZF99, FCW05, WCF07). Model~3 has the remarkably simple bivariate distribution $g(M, \\tau) \\propto M^{\\beta}\\tau^{\\gamma}$. The corresponding luminosity, mass, and age distributions are all pure power laws: $\\phi(L) \\propto L^{\\alpha}$, $\\psi(M) \\propto M^{\\beta}$, $\\chi(\\tau) \\propto \\tau^{\\gamma}$.  We have compared Models 1, 2, and 3 with the empirical distributions of luminosities, masses, and ages for a large sample of star clusters in the Antennae galaxies. These are based on our \\textit{UBVI}H$\\alpha$ photometry of point-like sources in images taken with the WFPC2 on {\\it HST\\/} and comparisons with stellar population models to estimate $M$ and $\\tau$ for each source (after correcting for interstellar reddening). To distinguish clusters of stars from individual stars, we have restricted our sample for analysis to objects brighter than $M_V = -9$, corresponding to $L > 3 \\times 10^5 L_{\\odot}$. This limit defines the domain of validity of the models in the $L$-$\\tau$ and $M$--$\\tau$ planes, the regions above the solid curves in Figures~1 and 2 respectively, corresponding approximately to $\\tau \\la 10^7 (M/10^4 M_{\\odot})^{1.3}$~yr. Thus, all the results presented in this paper pertain to relatively massive and relatively young clusters (except in \\S4.3).  We find that Model~3, with mass-{\\it in\\/}dependent disruption, provides a good match to the observed mass-age distribution of the Antennae clusters. The best-fitting exponents in this model are $\\beta \\approx -2$ and $\\gamma \\approx -1$. Models~1 and 2, with mass-dependent disruption fare much worse. Even with complete freedom to adjust several parameters, these models never come close to matching the data. While our detailed comparisons are based on models in which the formation rates of clusters are constant, our main conclusion, that the disruption rates do not depend on mass, is valid even if the formation rates are variable. As a check on these results, we have also rederived the luminosity function of the Antennae clusters. After corrections for interstellar extinction, we find a good fit to a pure power law, $\\phi(L) \\propto L^{\\alpha}$, with $\\alpha \\approx -2$. The fact that the mass and luminosity functions are nearly identical power laws ($\\alpha \\approx \\beta$) is a consequence of---and additional support for---the weak or nonexistent correlations between masses and ages, i.e., the decomposition $g(M,\\tau) \\propto \\psi(M)\\chi(\\tau)$. We have also investigated recent claims that there are bends or other features in the mass and luminosity functions of the Antennae clusters (Fritze-von Alvensleben 1999; Mengel et al. 2005; Anders et al. 2007). We show that these features are artifacts of the ways the samples are defined and do {\\it not\\/} reflect physical processes involved in the formation and disruption of the clusters.  In an effort to understand the physical basis for our simple---but doubtless approximate---model for $g(M,\\tau)$, we first note that the resulting mass function, $\\psi(M) \\propto M^{-2}$, resembles that of molecular clouds in nearby galaxies. This indicates that the efficiency of star formation in the clouds is roughly independent of their masses. We also consider a variety of processes that could disrupt the protoclusters and clusters: removal of ISM by stellar feedback, continued stellar mass loss, tidal disturbances by passing molecular clouds, gravitational shocks during rapid passages near the galactic bulge and through the galactic disk, orbital decay into the galactic center caused by dynamical friction, and stellar escape driven by internal two-body relaxation. We suggest that the massive young clusters in the Antennae galaxies are disrupted in the following approximate sequence: (1) ISM removal, $\\tau \\la 10^7$~yr, (2) stellar mass loss, $10^7 {\\rm yr} \\la \\tau \\la 10^8 {\\rm yr}$, (3) tidal disturbances, $\\tau \\ga 10^8$~yr. We have argued on theoretical grounds that these processes would operate at rates roughly independent of the masses of the clusters, consistent with our empirically-based decomposition $g(M, \\tau) \\propto \\psi(M)\\chi(\\tau)$ (which, however, is only well established for $\\tau \\la 10^8$~yr). In combination, these processes plausibly account for the observed decline of the age distribution, $\\chi(\\tau) \\propto \\tau^{-1}$. In the longer term, after 12 Gyr or so, the escape of stars driven by two-body relaxation will preferentially disrupt low-mass clusters and imprint a peak or turnover in the evolved mass function at $M_p \\approx (1-2) \\times 10^5 M_{\\odot}$, similar to that for old globular clusters in the Milky Way and other galaxies.  How general is this picture? In the Introduction, we mentioned some circumstantial evidence in favor of its wide applicability: the statistics of embedded clusters in the solar neighborhood (Lada \\& Lada 2003), the luminosities of the brightest clusters in different galaxies (Larsen 2002; Whitmore 2003; WCF07), the similarity of the mass spectra of molecular clouds in different galaxies (Blitz et al. 2007), and the fact that the early disruption of protoclusters and clusters ($\\tau \\la 10^8$~yr) is driven mainly by internal processes that depend weakly, if at all, on the properties of their host galaxies (\\S4 here). Furthermore, from a detailed study of the mass-age distributions of clusters in the LMC and SMC, we find results nearly identical to those presented here for the clusters in the Antennae (CFW09). This is a strong test because the Antennae and the Magellanic Clouds represent very different environments for star and cluster formation: two large interacting galaxies on one hand, two small, relatively quiescent galaxies on the other hand. Nevertheless, further tests of this picture in other galaxies would be beneficial. If this picture turns out to be generally valid, it will mean that the main difference between populations of young clusters is simply in the normalization of $g(M, \\tau)$ and hence in the overall formation rate. From this perspective, the objects traditionally designated open, populous, globular, or super clusters would belong to a continuum, with mass and age as the most fundamental variables, and would not require different formation and/or disruption processes to account for their observed properties.  The picture presented here has the virtues of simplicity and possible universality. It captures the salient properties of the observed mass-age distribution, at least for the populations of clusters we have studied so far. Nevertheless, we do not expect our model $g(M, \\tau)$ to be strictly universal. It may break down outside its domain of validity in most galaxies or inside this domain in some extreme environments, such as those dominated by population III stars. There may also be minor deviations from the model in some normal galaxies, even within the domain of validity, caused for example by variations in the formation rates and/or differences in the disruption rates by passing molecular clouds. However, these are only likely to affect $g(M, \\tau)$ at intermediate ages ($10^8 {\\rm yr} \\la \\tau \\la 10^9{\\rm yr}$), which will be difficult to detect observationally. We see an analogy between, on one hand, the potentially universal stellar IMF, represented by a simple, approximate formula (first by Salpeter (1955) and then revised in subsequent studies based on more modern data) and, on the other hand, the potentially universal mass-age distribution $g(M, \\tau)$ for star clusters, represented here by a simple, approximate formula. Both the stellar IMF and the cluster $g(M, \\tau)$ are thought to be the outcome of several complex physical processes, in ways not yet fully understood, and neither is expected to hold exactly in all situations. Even so, these functions provide compact and useful summaries of the observations and are thus natural focal points for future studies of the formation and early evolution of stars and clusters."
    },
    "0910/0910.3277_arXiv.txt": {
        "abstract": "Imaging atmospheric Cherenkov telescopes (IACTs) used for ground-based gamma-ray astronomy at TeV energies use reflectors with areas on the order of 100m$^2$ as their primary optic.  These tessellated reflectors comprise hundreds of mirror facets mounted on a space frame to achieve this large area at a reasonable cost.  To achieve a reflecting surface of sufficient quality one must precisely orient each facet using a procedure known as alignment.  We describe here an alignment system which uses a digital (CCD) camera placed at the focus of the optical system, facing the reflector.  The camera acquires a series of images of the reflector while the telescope scans a grid of points centred on the direction of a bright star.  Correctly aligned facets are brightest when the telescope is pointed directly at the star, while mis-aligned facets are brightest when the angle between the star and the telescope pointing direction is twice the misalignment angle of the facet.  Data from this scan can be used to calculate the adjustments required to align each facet.  We have constructed such a system and have tested it on three of the VERITAS IACTs.  Using this system the optical point spread functions of the telescopes have been narrowed by more than 30\\%.  We present here a description of the system and results from initial use. ",
        "introduction": "The current generation of imaging atmospheric Cherenkov telescopes operating around the world~\\cite{Holder08,Hinton04,Baixeras04} has ushered in a new era in TeV gamma-ray astronomy.  The number of detected TeV gamma-ray sources has grown from below ten in 2000 to more than seventy today~\\cite{Aharonian08} largely because of the increased sensitivity of the instrumentation.  This increase results from the use of the following: \\begin{itemize} \\item{larger reflectors} \\item{cameras with larger fields of view and higher resolution} \\item{multiple telescopes making stereoscopic observations} \\item{flash-ADC-based data acquisition systems} \\end{itemize}  For the benefits afforded by large reflectors and high resolution cameras to be fully realised, the optical quality of the telescopes must be maintained at a high level.  Since the reflectors of these telescopes comprise several hundred mirror facets, their alignment presents a significant technical and logistical challenge.  The VERITAS array, located in southern Arizona, USA, employs four twelve-metre-diameter f-1.0 reflectors of the Davies-Cotton type~\\cite{DaviesCotton}.  Each reflector consists of a tubular steel optical support structure (OSS) on which 345 identical hexagonal mirror facets are mounted.  The facet mounts allow precision adjustments to bring the focus of each to the same point on the primary focal plane of the telescope.  Since the mirror facets are exposed to the dust of the Arizona Sonoran desert their reflectivity degrades over time so they are therefore re-coated on a regular basis~\\cite{Roache08}.  This process maintains the reflectivity of the facets but their removal and re-installation compromises the optical quality of the reflector as a whole so the alignment of the facets must be repeated on a regular basis.  We present here an alignment system, based on the technique originally suggested by Arqueros et al.~\\cite{Arqueros05}, which can be used for aligning the VERITAS telescopes.  It achieves the quality desired in a reasonable length of time at modest cost.  Importantly, the optimal alignment is achieved for typical observation elevations. ",
        "conclusions": "An alignment system based on the technique suggested by \\cite{Arqueros05} has been developed and used to improve the optics of three VERITAS telescopes.  This has led to a reduction in the size of the PSF by more than 30\\%.  Moreover this system is less labour-intensive than that which was previously used.  It has the advantage that the telescope reflectors are directly optimised for use at typical observing elevations.  Further investigations are planned."
    },
    "0910/0910.2792_arXiv.txt": {
        "abstract": "We present the detailed optical to far-infrared observations of SST J1604+4304, an ULIRG at $z = 1.135$. Analyzing the stellar absorption lines, namely, the CaII H \\& K and Balmer H lines in the optical spectrum, we derive the upper limits of an age for the stellar population. Given this constraint, the minimum $\\chi^2$ method is used to fit the stellar population models to the observed SED from 0.44 to 5.8\\micron. We find the following properties. The stellar population has an age 40 - 200 Myr with a metallicity 2.5 $Z_{\\sun}$. The starlight is reddened by $E(B-V) = 0.8$. The reddening is caused by the foreground dust screen, indicating that dust is depleted in the starburst site and the starburst site is surrounded by a dust shell. The infrared (8-1000\\micron) luminosity is $L_{ir} = 1.78 \\pm 0.63 \\times 10^{12} L_{\\sun}$. This is two times greater than that expected from the observed starlight, suggesting either that 1/2 of the starburst site is completely obscured at UV-optical wavelengths, or that 1/2 of $L_{ir}$ comes from AGN emission. The inferred dust mass is $2.0 \\pm 1.0 \\times 10^8 M_{\\sun}$. This is sufficient to form a shell surrounding the galaxy with an optical depth $E(B-V) = 0.8$. From our best stellar population model - an instantaneous starburst with an age 40 Myr, we infer the rate of 19 supernovae(SNe) per year. Simply analytical models imply that 2.5 $Z_{\\sun}$ in stars was reached when the gas mass reduced to 30\\% of the galaxy mass. The gas metallcity is $4.8 Z_{\\sun}$ at this point. The gas-to-dust mass ratio is then $120 \\pm 73$. The inferred dust production rate is $0.24 \\pm 0.12 M_{\\sun}$ per SN. If  1/2 of $L_{ir}$ comes from AGN emission, the rate is $0.48 \\pm 0.24 M_{\\sun}$ per SN. We discuss the evolutionary link of SST J1604+4304 to other galaxy populations in terms of the stellar masses and the galactic winds, including optically selected low-luminosity Lyman $\\alpha$-emitters and submillimeter selected high-luminosity galaxies. ",
        "introduction": "$IRAS$(Infrared Astronomical Satellite) observations discovered numerous infrared galaxies. At bolometric luminosities $> 10^{11} L_{\\sun}$, infrared galaxies are the dominant population in the local Universe, being more numerous than optically selected starbursts, AGNs, and quasars \\citep{sanders96}. By luminosity, infrared galaxies are classified into luminous infrared galaxies(LIRGs) with $L_{ir} > 10^{11} L_{\\sun}$, ultraluminous infrared galaxies(ULIRGs) with $L_{ir}> 10^{12} L_{\\sun}$, and hyperluminous infrared galaxies(HyLIRGs) with $L_{ir}> 10^{13} L_{\\sun}$. The ratio of infrared to visible luminosity, $L_{ir}/L_B$ increases with $L_{ir}$ \\citep{soifer89}. Although there is evidence that an optically buried AGN may exist in LIRGs, there is similarly strong evidence that enhanced star formation is ongoing (\\citealt{armus95}; \\citealt{sanders96}). \\citet{heckman90} found that their optical spectroscopic data support the superwind model in which the kinetic energy provided by supernovae (SNe) and winds from massive stars in starburst site drives a large-scale outflow.  The {\\it Infrared Space Observatory (ISO)} source counts in the mid- and far-infrared extended our knowledge of infrared galaxies up to $z \\sim 1$ (\\citealt{taniguchi97};\\citealt{oliver97};\\citealt{kawara98}; \\citealt{puget99}; \\citealt{elbaz99};\\citealt{serjeant00}; \\citealt{sato03}; \\citealt{kawara04}),and revealed the number of ULIRGs rapidly grew with increasing redshift(\\citealt{genzel00}; \\citealt{chary01}; \\citealt{heraudeau04}; \\citealt{oyabu05}).  The next large advance was made by the {\\it Spitzer Space Telescope} with the sensitivity to probe infrared galaxies at $z = 2- 4$ (see review in \\citealt{soifer08} reference therein). \\citet{reddy06}, using deep {\\it Spitzer} 24\\micron\\ observations, examined star formation and extinction in optically selected galaxies, near-infrared-selected galaxies, and submillimeter-selected galaxies (SMGs) at $z \\sim 2$. {\\it Spitzer} provided a powerful probe to measure redshifts of SMGs, the most luminous infrared galaxy population as HyLIRGs \\citep{pope08a}. Based on {\\it Spitzer} mid-infrared data, the diagnostic diagram to separate starburst- and AGN-dominated infrared galaxies has been developed(e.g., \\citealt{stern05}; \\citealt{pope08b}; \\citealt{polletta08}), and it is suggested that $L_{ir}$ in SMGs as well as ULIRGs at $z \\sim 2$ is dominated by star formation\\citep{pope08a}, while many SMGs contain an AGN as evidenced by their X-ray properties. \\citet{dey08} uncovered a new population of dust-obscured galaxies (DOGs). These galaxies have $L_{ir}$ comparable to ULIRGs. However, their $L_{ir}/L_B$ is greater than local ULIRGs. DOGs and SMGs as well as high-redshift ($z \\sim 1$) ULIRGs are generally considered to be young and massive galaxies, although the nature of the stellar populations is far from being understood. To understand the evolutionary link of these infrared galaxies to optically selected and near-infrared-selected galaxies, it is crucial to observe stellar absorption lines which indicate the age of the stellar populations.  During the course of searching for distant objects in a cluster field, SST J1604+4304 drew our attention with its red optical-3.6\\micron\\ color. Follow-up study revealed this object is an ULIRG at $z \\sim 1.135$ and its energy source is young stars with an age $< 200$ Myr.  We present the data consisting of optical spectroscopy and photometry from the optical to far-infrared in section 2. We derive an age and metallicity of the stellar population, and show a foreground dust screen is plausible in Section 3. In Section 4, mid- and far-infrared emission by dust and energetics are discussed. In Section 5, we discuss the inferred size of a dust shell surrounding the galaxy, the SN rate and metallicity expected from the broad-band SED analysis, and the dust production rate per SN. The evolutionary link of infrared galaxies to optically selected galaxies is discussed based on the galactic wind models of elliptical galaxies.  We adopt $H_0$ = 70 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_m = 0.3$, and $\\Omega_{\\Lambda} = 1 - \\Omega_m$ throughout this paper. In this cosmology, the luminosity distance to the object is 7520 Mpc. The flux density $F_{\\nu}$ is used in units of $\\mu$Jy, which is converted to AB magnitudes through the relation  $m(AB) = -2.5 log(F_{\\nu})$ + 23.9.   \\section[]{Observations and data analysis}  This work is based on the archival data taken by the {\\it Spitzer Space Telescope}, the {\\it Hubble Space Telescope (HST)}, and the $Subaru$ Telescope and new photometric observations performed on the {\\it United Kingdom Infrared Telescope(UKIRT)} and the $MAGNUM$\\footnote{This telescope is dedicated to the Multicolor Active Galactic Nuclei Monitoring (MAGNUM) project and is installed at the Haleakala Observatories in Hawaii \\citep{yoshii02}.} telescope. In addition, the {\\it Gemini Telescope North} was used to do optical spectroscopy.   \\begin{figure} \\includegraphics[width=84mm]{k3_chart.ps} \\caption{Images centered at SST J1604+4304, where the size is 70\\arcsec $\\times$ 70\\arcsec for $I_{814}$, 3.6\\micron, 8.0\\micron, and 24\\micron, and 240\\arcsec $\\times$ 240\\arcsec for 70\\micron, and 160\\micron. For all plots, north is at the top and east is to the left. We present 32\\arcsec radius circles which are the size of the aperture used for MIPS 160\\micron\\ photometry. Two bars in the $I_{814}$ image are used to indicate SST J1604+4304. } \\label{f_chart} \\end{figure}  \\begin{figure} \\includegraphics[width=84mm]{k3_closeup.ps} \\caption{10\\arcsec $\\times$ 10\\arcsec images centered at SST J1604+4304. For all plots, north is at the top and east is to the left. We present 3\\farcs7 radius circles which are the size of the aperture used for IRAC 3.6-8.0\\micron\\ photometry. } \\label{f_closeup} \\end{figure}  \\subsection{Target}  SST J1604+4304 has been found at \\\\ \\begin{flushleft} R.A. 16$^h$ 04$^m$ 25\\fs538\\\\ decl. +43\\degr 04\\arcmin 26\\farcs55,\\\\ \\end{flushleft} \\noindent in the J2000 system in a field observed with the $Spitzer$ IRAC instrument. The optical counterpart in the $HST$ data is located within the 0\\farcs4 offset accuracy requirements of $Spitzer$ \\citep{werner04}. Images of SST J1604+4304 are given in Figures \\ref{f_chart} and \\ref{f_closeup}. The broad-band fluxes are given in Table \\ref{t_fluxes}.  SST J1604+4304 is only 32\\arcsec away from the center of the massive galaxy cluster CL 1604+4304 at $z = 0.90$ \\citep{gal04}. Based on a weak-lensing analysis, \\citet{margoniner05} find the mass contained within projected radius $R$ is $(3.69 \\pm 1.47)[R/(500 kpc)] \\times 10^{14} M_{\\sun}$, corresponding to an inferred velocity dispersion 1004 $\\pm$ 199 km s$^{-1}$. Thus, using a singular isothermal sphere, we obtain the magnification $m = 1.17 \\pm 0.09$ for SST J1604+4304 \\citep{schneider92}.   \\subsection{$Spitzer$ data}  All IRAC and MIPS data sets were retrieved from the $Spitzer$ public archive. The IRAC data consist of four broadband images at 3.6, 4.5, 5.8, and 8.0 \\micron. The field-of-view is nearly 5\\farcm2 $\\times$ 5\\farcm2 with a linear scale of approximately 1\\farcs22 pixel$^{-1}$. The reader is referred to \\citet{fazio04} for a complete report on the in-flight performance of IRAC. In the pipeline processing, ten BCD FITS images were co-added into a single mosaic image per band with a technique similar to 'drizzle'. The total effective exposure time is 1936 sec pixel$^{-1}$ per band. Photometry was performed using a 3 pixel (3\\farcs7) radius aperture. As can be seen in the 3.6\\micron\\ image in Figure \\ref{f_chart}, there are two sources which can contribute to the flux within the photometric aperture; a faint source is at 6\\farcs5 to the southwest and a bright source 10\\arcsec to the west. Using the point-spread function which is available on the IRAC website, their contributions were estimated to be 3.1 - 4.3 \\% of the flux of SST J1604+4304 depending on the photometric band. The total fluxes after applying the aperture correction are given in Table \\ref{t_fluxes}. The quoted errors include 5\\% absolute calibration uncertainty that SSC\\footnote{See Infrared Array Camera Data Handbook} recommends for general observers. Note that the color correction is very small and ignored.   The MIPS data were taken at 24, 70, and 160 \\micron\\ in 2004 March. In the pipeline processing, the mapped BCD data were rebinned and coadded to a single image with a linear pixel scale of 2\\farcs45, 4\\farcs0, and 4\\farcs0 pixel$^{-1}$ for 24, 70, and 140 \\micron, respectively. The total integration times are 93, 53, and 84 sec pixel$^{-1}$. In the following analysis, we use the following post-BCD products, namely, mosaic images for 24\\micron\\ and filtered mosaic images for 70 and 160\\micron. 24\\micron\\ photometry was performed using a 7\\arcsec radius aperture. The contributions, to the flux within the photometric aperture, from the two sources (see the 24\\micron\\ image in Figure \\ref{f_chart}), located approximately at 15\\arcsec to the west-northwest and 13\\arcsec to the southwest, were estimated using the point-spread function on the MIPS website. Their contributions are 9\\% of the flux within the aperture. Then, the color correction for 500K blackbody and the aperture correction were applied.  MIPS 70 and 160 \\micron\\ photometry was performed using 16\\arcsec and 32\\arcsec radius apertures, respectively. As seen in Figure \\ref{f_chart}, SST J1604+4304 is faint in the far-infrared. We measured sky fluxes with the same aperture at 10 or 11 positions along the annulus with radii of 36\\arcsec - 72\\arcsec at 70\\micron\\ and 96 - 150\\arcsec at 160\\micron. The flux averaged over these sky fluxes was subtracted from the flux measured within the object aperture centered at SST J1604+4304. The aperture correction and the color correction for 30K were applied. The detections are 2.5$\\sigma$ and 5.0$\\sigma$ at 70 and 160\\micron, respectively. As seen in the 24\\micron\\ image in Figure \\ref{f_chart}, there are several 24\\micron\\ sources within the 160\\micron\\ photometry aperture. Judging from their [3.6\\micron]-[24\\micron] color and 70 and 160 \\micron\\ brightness at their respective positions, the source at 13\\farcs5 to the west-northwest, the only source as red as SST J1604+4304, may contribute 1/3 of the flux within the aperture at most. The 160 \\micron\\ flux is 38.4 $\\pm$ 7.65 mJy if the contributions from this source can be ignored, while the flux is 25.6 $\\pm$ 7.65 mJy if this source contributes 1/3 of the flux within the aperture. The real flux of SST J1604+4304 should be between them, and thus the flux is 32 $\\pm$ 14 mJy. The aperture correction and the color correction for 30K blackbody were applied. Note that the 24\\micron\\ error in Table \\ref{t_fluxes} includes 4\\% absolute calibration uncertainty\\footnote{Multiband Imaging Photometer for Spitzer (MIPS) Data Handbook}.   \\subsection{Optical imaging data}  The $V_{606}$, $I_{814}$, and $Z_{850}$ image data were retrieved from the $HST$ public archive, where $V_{606}$, $I_{814}$, and $Z_{850}$ denote the $F606W$, $F814W$, and $F850LP$, respectively. The Advanced Camera for Surveys (ACS) Wide Field Camera (WFC) was used, which covers a field of 3\\farcm4 $\\times$ 3\\farcm4 with a scale of 0\\farcs05 pixel$^{-1}$. Four image data were combined into a single co-added image per band. The total exposure times are 4840 sec pixel$^{-1}$ for $V_{606}$ and $I_{814}$ and 6000 sec pixel$^{-1}$ for $Z_{850}$. The $I_{814}$ image in Figure \\ref{f_closeup} shows that SST J1604+4304 has irregular morphology extending 1\\farcs6 in the NW direction with a 0\\farcs6 width. To determine the total flux of this object, $I_{814}$ photometry was performed using 0\\farcs8, 1\\farcs0, 1\\farcs2, and 2\\farcs0 radius apertures. The $I_{814}$ flux grows from 0\\farcs8 to 1\\farcs2, while it barely increase from 1\\farcs2 to 2\\farcs0. Note that contributions from the two nearest sources were subtracted from the flux measured with the 2\\farcs0 radius aperture. Thus, the flux measured with the 1\\farcs2 radius aperture represents the total flux of SST J1604+4304. The ratio of the flux with a 0\\farcs8 aperture to a 1\\farcs2 aperture is 1.24 $\\pm$ 0.05. $V_{606}$ and $Z_{850}$ photometry was performed using a 0\\farcs8 aperture. The resultant fluxes were converted to the total fluxes by multiplying 1.24 $\\pm$ 0.05, because we did not measure significant changes in optical colors across the target. The total fluxes are given in Table \\ref{t_fluxes}. The errors in Table \\ref{t_fluxes} include the error of aperture conversion from 0\\farcs8 - 1\\farcs2 and the correction for correlated noise in drizzling.   $Subaru$ $B, V, R_c, I_c, z$ data were retrieved from the SMOKA science achieve. The data were taken with Suprime-Cam (SUP) imager which has a field-of-view of 34\\arcmin $\\times$ 27\\arcmin with a scale of 0\\farcs2 pixel$^{-1}$. The total exposure time is 2880, 2520, 3600, 1680, 3300 sec pixel$^{-1}$ after co-adding 4 to 11 image frames, using the standard SUP script \\citep{yagi02}. Flux scaling was performed observing 11 standard stars given by \\citet{majewski94}. $I_c$ photometry were performed using three apertures with a radius of 1\\farcs0, 1\\farcs2, and 2\\farcs0, and it was confirmed the $Subaru$ $I_c$ brightness distribution is identical to that measured at the $HST$ $I_{814}$ photometry band. Thus, we use the 1\\farcs2 radius aperture centered on SST J1604+4304 to obtain the fluxes given in Table \\ref{t_fluxes}. The errors in Table \\ref{t_fluxes} include 5\\% photometric uncertainty.   \\subsection{H/K imaging observations}  The $H$-band image was taken with the UIST (1-5 micron imager spectrometer) on $UKIRT$ in February 2006. UIST was used in the imaging mode with a field-of-view of 2\\arcmin $\\times$ 2\\arcmin and a scale of 0\\farcs12 pixel$^{-1}$. The total exposure time is 1800 sec pixel$^{-1}$ after co-adding individual image data with a integration time of 60 sec. The sky condition was photometric with 0\\farcs6 seeing. The K-band image was taken with the multicolor imaging photometer (MIP) on the 2 m $MAGNUM$ telescope (\\citealt{kobayashi98};\\citealt{minezaki04}) in August 2006. MIP has a field-of-view of 1\\farcm5 $\\times$ 1\\farcm5 with a linear scale of 0\\farcs346 pixel$^{-1}$. The total integration time is 2340 sec pixel$^{-1}$ after co-adding individual image data. The weather condition was photometric with 1\\arcsec seeing. Flux calibration was performed using 2MASS 16042631+4303413 (H=14.36, K= 14.39) observed simultaneously with SST J1604+4304. Photometric uncertainty is 5\\% and 9\\% for H and K, respectively.  \\begin{table} \\centering \\caption{Multiwavelength photometry of SST J1604+4304.} \\begin{tabular}{crrrr} \\hline Filter$^a$   &  Total flux        & $R_{ph}$$^b$  &  Instrument & UT \\\\ band     &  $\\mu$Jy              & arcsec     &             & yy/mm    \\\\ \\hline $V_{606}$  & 0.224 $\\pm$ 0.062 & 1.2      & $HST$ ACS   & 03/08 \\\\ $I_{814}$  & $1.05 \\pm 0.08$   & 1.2      & $HST$ ACS   & 03/08 \\\\ $Z_{850}$  & $1.53 \\pm 0.13$   & 1.2      & $HST$ ACS   & 07/02 \\\\ 3.6\\micron    & $60.35 \\pm 3.16$  & 3.7      & $Spitzer$ IRAC & 04/03 \\\\ 4.5\\micron    & $53.76 \\pm 2.79$  & 3.7      & $Spitzer$ IRAC & 04/03 \\\\ 5.8\\micron    & $55.11 \\pm 4.07$  & 3.7      & $Spitzer$ IRAC & 04/03 \\\\ 8.0\\micron & $65.68 \\pm 5.36$     & 3.7      & $Spitzer$ IRAC & 04/03 \\\\ 24\\micron  & $319 \\pm 18.4$       & 7.0      & $Spitzer$ MIPS & 04/03 \\\\ 70\\micron  & $2150 \\pm 870$       & 16      & $Spitzer$ MIPS & 04/03 \\\\ 160\\micron & $32000 \\pm 14000$     & 32      & $Spitzer$ MIPS & 04/03 \\\\ \\hline $B$     & $0.0803 \\pm 0.013$ & 1.2 & $Subaru$ SUP & 00/06 \\\\ $V$     & $0.111 \\pm 0.015$  & 1.2 & $Subaru$ SUP & 05/05 \\\\ $R_c$   & $0.218 \\pm 0.018$  & 1.2 & $Subaru$ SUP & 01/04,05 \\\\ $I_c$   & $0.689 \\pm 0.042$  & 1.2 & $Subaru$ SUP & 01/04 \\\\ $z$     & $2.29 \\pm 0.21$    & 1.2 & $Subaru$ SUP & 01/06 \\\\ $H$     & $5.40 \\pm 1.19$    & 1.2 & $UKIRT$  UIST & 06/02 \\\\ $K$     & $20.06 \\pm 3.35$   & 1.2 & $MAGNUM$ MIP  & 06/08 \\\\ \\hline \\hline 20cm    & $< 0.66$ mJy$^c$       & 2.7 & $VLA$$^e$           & 1995 \\\\ 0.2-2 keV & $<2.1$ $^d$ & 15 & $XMM$$^f$ & 02/02 \\\\ 2-10 keV & $<15$ $^d$ & 15 & $XMM$$^f$ & 02/02 \\\\ \\hline \\end{tabular} \\medskip \\\\ \\begin{flushleft} $^a$ $V_{606}, I_{814}, \\& Z_{850}$ denote F606W,F814W,\\& F850LP, respectively.\\\\ $^b$ The radius of the photometric apertures which were used to obtain the total fluxes.\\\\ $^c$ To 3 $\\sigma$.\\\\ $^d$ To 3 $\\sigma$ in units of $10^{-15}$ erg cm$^{-2}$ s$^{-1}$.\\\\ $^e$ From the FIRST survey \\citep{white97}.\\\\ $^f$ Observed by \\citet{lubin04}. The FLIX online server was used to derive the upper limit. \\end{flushleft} \\label{t_fluxes} \\end{table}  \\subsection{Optical spectroscopy}  Optical spectroscopy was carried out using the Gemini Multi-Object Spectrograph (GMOS) on the Gemini-North telescope in 2007 August. The R150\\_G5306 grating was used along with the OG151\\_G036 filter, covering 5,000 - 10,000 \\AA\\ simultaneously with 6.96 \\AA/pixel after binned by 4 in the dispersion direction. The different wavelength centers (8,200 and 8,250 \\AA) were set to fill the CCD gaps. A 1\\farcs0 width slit was set so that the dispersion direction is perpendicular to the long axis of the target. Seven 1800 sec exposures were taken in the Nod \\& Shuffle mode. Data were reduced in a standard manner using the IRAF GEMINI.GMOS package. The sky subtraction was made with the GNSSKYSUB routine. Wavelengths were scaled with CuAr spectra and the relative photometric calibration was performed using the standard star (HZ44). The results are shown in the bottom panel of Figure \\ref{f_sed}. The spectrum which is median-smoothed with a width of 40 pixles (i.e., $\\sim$ 280 \\AA), is compared with the photometry from the images in the top panel of Figure \\ref{f_sed}. The two measurements agree with each other.  \\begin{figure} \\includegraphics[width=84mm]{k3_sed.ps} \\caption{The top panel shows UV to far-infrared SED for SST J1604+4304. The broad-band data are shown along with the median-smoothed optical spectrum ({\\it gray}) taken with Gemini-North GMOS. The bottom panel shows the optical spectrum ({\\it top}) together with the synthetic spectrum ({\\it bottom}) at an age of 400 Myr, reddened by E(B-V) = 0.6} \\label{f_sed} \\end{figure} ",
        "conclusions": "For simplicity,  we temporarily assume that SST J1604+4304 is purely powered by star formation and at the end we will discuss the impact if SST J1604+4304 is AGN-dominated.  \\subsection{Stellar mass}  The stellar mass $M_*$ can be estimated as $L_{ir}/(L/M_*)$, where $L/M_*$ is the luminosity to mass ratio for the stellar population models. Thus, we have $M_* = 6 - 13 \\times 10^{10} M_{\\sun}$ at $ t = 0$ as listed in Table \\ref{t_properties}. If this stellar mass was constantly built up during 40 - 200 Myr, the star formation rate is $SFR$ = 600 - 1650 M$_{\\sun}$ yr$^{-1}$. Using the calibrations of SFR in terms of the [OII] luminosity (\\citealt{kennicutt98}), after dereddening the [OII] flux, implies 13 - 500 $M_{\\sun}$ yr$^{-1}$. Considering that we are only looking at 1/2 of the starburst activity in the the optical, we obtain SFR = 26 - 1000 $M_{\\sun}$ yr$^{-1}$, in good agreement with the previous estimate.  \\subsection{Comparison with other infrared galaxy populations}  SST J1604+4304 is a dusty, super metal-rich, young galaxy at $z = 1.135$ with $E_{B-V}$ = 0.8, metallicity 2.5 $Z_{\\sun}$, and an age 40 - 200 Myr. The age  is comparable to young, optical-selected populations like Lyman-$\\alpha$ emitters (LAEs) and Lyman break galaxies (LBGs). Studying the $z \\sim 3$ LAE sample, \\citet{gawiser06} inferred ages of the order of 100 Myr with stellar masses $5 \\times 10^8 M_{\\sun}$ and no dust extinction. \\citet{pirzkal07} showed $4 < z < 5.7$ LAEs had very young ages a few Mys with stellar masses $10^6 - 10^8 M_{\\sun}$. \\citet{sawicki98}, studying $z > 2$ LBGs, found that their LBGs were dominated by young stellar populations with ages $<$ 200 Myr with stellar masses several $\\times 10^9 M_{\\sun}$ and moderate dust extinction typically $E(B-V) \\sim 0.3$. The $z = 2 - 3.5$ sample of LBGs by \\citet{papovich01} has ages of 30 Myr to 1 Gyr with stellar masses $10^9 - 10^{11} M_{\\sun}$ with $E(B-V)$ = 0.0 - 0.4. SST J1604+4304, having a stellar mass 6 - 13 $\\times 10^{10} M_{\\sun}$, is more massive and more dusty than these optically selected galaxies.  SST J1604+4304 having $f_{\\nu}(24\\mu m)/f_{\\nu}(R) > 1000$ meets the criteria for 24\\micron-selected high-redshift ($z \\sim 2$) dust-obscured galaxies (DOGs). DOGs and submillimeter-selected galaxies (SMGs) have extremely large $L_{ir}/L_B$ which cannot be accounted for by redshifting local ULIRGs (\\citealt{dey08};\\citealt{pope08b}). \\citet{dey08} found nearly all of DOGs are ULIRGs with $L_{IR} > 10^{12} L_{\\sun}$ and more than half of these have $L_{IR} > 10^{13} L_{\\sun}$. The infrared luminosity of radio-identified SMGs ranges $L_{IR} = 2 \\times 10^{11} - 10^{14} L_{\\sun}$ \\citep{chapman05}. A SFR 1000 $M_{\\sun} yr^{-1}$, which corresponds to $L_{IR} = 5 \\times 10^{12} L_{\\sun}$ \\citep{kennicutt98}, builds up a stellar mass of $10^{11} M_{\\sun}$ in 100 Myr. Thus, it is considered that DOGs and SMGs are progenitors of present-day massive galaxies. If ages of stellar populations in DOGs and SMGs are derived, as we did for SST J1604+4304 by measuring the stellar absorption lines, the evolutionary link of DOGs and SMGs to normal ULIRGs and optically selected galaxies will be better understood.  \\subsection{SNe for dust and metal production}  In SST J1604+4304 with an age 40 - 200 Myr, most of the dust would be formed by type II supernovae (SNe II), because low-mass stars had not evolved to form dust in the expanding envelope. This may be the reason why our best model results in poor goodness of fit $\\chi_{\\nu}^2(min) = 5.2$. The Calzetti's extinction law may not be applicable to such young galaxies, because dust grains in local starbursts would be dominated by those formed in the outflows of low-mass stars and not by SNe II dust. In this regard, it is of great interest to use the SED of SST J1604+4304 to test extinction laws such as those calculated by Hirashita et al. (2005, 2008) that are predicted for dust formed in SNe II ejecta \\citep{nozawa03}.  Our best model is the instantaneous-burst model with an age 40 Myr with a stellar mass $6.4 \\pm 2.3 \\times 10^{10} M_{\\sun}$ at $t = 0$. The metallicity is 2.5 $Z_{\\sun}$ and the dust mass is $M_{dust} = 2.0 \\pm 1.0 \\times 10^8 M_{\\sun}$. Until 40 Myrs after the onset of the starburst, stars with a mass $> 8 M_{\\sun}$ died with a SN explosion. The number of stars evolved into SNe in 40 Myr is $7.6 \\times 10^8$ or 19 SNe per year. Adopting the nucleosynthesis models by \\citet{nomoto06}, the total yield is 0.09 - 0.11 and the metal yield is 0.016 - 0.020 for the Chabrier IMF.\\footnote{The original yields are given for the Salpeter IMF. These are converted to the Chabrier IMF by simply multiplying 1.64 to adjust the difference in the mass fraction of stars from 8 - 100 $M_{\\sun}$ to those from 0.1 -100 $M_{\\sun}$.} The gas and stetllar metallicities can be estimated applying simply analytical models based on the instantaneous recycling approximation. The gas metallicity is expressed as $Z_g(t)$ = $y(t)ln(1/f(t))$ where $f(t)$ is the ratio of the gas mass to the galaxy mass, and $y(t)$ is the ratio of the rate at which the metal is produced by the events of nucleosysnthesis and ejected into the interstellar gas to the rate at which hydrogen is removed from the interstellar gas by star formation \\citep{searle72}. The stellar metallicity $Z_*$ is derived from $S(Z)/S_1 = (1 - f_1^{Z/Z_1})/(1 - f_1)$ and $Z_* = \\int^{Z_1}_0 ZdS(Z)/ \\int^{Z_1}_0 dS(Z)$, where $S(Z)$ is the total mass of stars born up to a time when the metal abundance reached to the value $Z$, and $S_1, Z_1, f_1$ denote $S(t_1), Z(t_1), f(t_1)$, respectively \\citep{pagel75}. In our case, $t_1$ is a time when the stellar metallicity reached to $Z_* = 2.5 Z_{\\sun}$. When the gas mass is reduced to $ f = 0.3$, we have $Z_* = 2.4 - 2.7 Z_{\\sun}$ and $Z_g = 4.5 - 5.1 Z_{\\sun} $. The inferred galaxy mass is $1.2 \\times$ the stellar mass at $t =0$, because the total mass ejected from stars at the events of SN explosions into the interstellar gas is 0.1 per unit stellar mass at $t =0$. The total gas mass is $2.3 \\times 10^{10} M_{\\sun}$. The gas-to-dust mass ratio is then $120 \\pm 73$, which is comparable to the value obtained in a $z = 1.3$ HyLIRG ($198 \\pm 53$) by \\citet{iono06}.  The dust production rate is $0.24 \\pm 0.12 M_{\\sun}$ per SN. This is consistent with the numerical results of dust formation in SNe; for progenitor masses ranging 13 - 30 $M_{\\sun}$, 2 - 5\\% of the progenitor mass is locked into dust grains at SN explosions \\citep{nozawa03}, 20 - 100\\% of which is destroyed by processing through the collisions with the reverse shocks resulting from the interaction of SN ejecta and with the ambient medium \\citep{nozawa07}.  \\subsection{Foreground dust screen}  As demonstrated in Section 3.2, a foreground dust screen geometry is plausible for SST J1604+4304. This indicates that dust is depleted in the starburst site and the stellar spectra are reddened by the foreground dust screen which enclosed the starburst site. Let's suppose a spherical shell with a radius $l$ for the dust distribution, then we observe $f_{\\nu} = 4 \\tau_{\\nu}B_{\\nu}(T_{dust})\\pi(l/D)^2$ in the rest-frame, where the thickness of the shell is $\\tau_{\\nu}$ which is 0.0074 at rest 80\\micron\\ corresponding to $E(B-V)$ = 0.83 with $R_V$ = 4.05 \\citep{draine03}, $D$ is the distance from the observer to SST J1604+4304, and $T_{dust} = 32.5 - 35 K$. Converting this relation into the observer's frame, we obtain the radius $l$ = 4.5 - 5.5 kpc or 0.56 - 0.65\\arcsec. This is comparable to the optical size of SST J1604+4304 which is approximately 1.2\\arcsec $\\times$ 0.5\\arcsec in the $I_{814}$ band. This supports the view that the galaxy is surrounded by the shell. Further support is the fact that there are no significant color gradients across the galaxy in the $HST$ images.  \\subsection{Evolutionary link to other galaxy populations}  As discussed by \\citet{calzetti01}, starburst environments are rather inhospitable to dust; dust grains in the starburst site can be transported to a large distance in a relatively short time by radiation pressure(\\citealt{ferrara91}; \\citealt{venkatesan06}), as well dust grains that formed at SN explosions are processed and evolve in SN remnants \\citep{nozawa07} - small size grains are quickly destroyed in SN remnants by sputtering. Thus, a cavity-shell structure is a natural geometry for the dust distribution in star-forming galaxies; the starburst site is inside the cavity where dust is depleted, and the opaque dust shell is surrounding the cavity. The shell is observed as a foreground screen. Such foreground screens were found in local starburst galaxies (\\citealt{calzetti94}; \\citealt{meurer95}). In the foreground screen of SST J1604+4304, the dust would be unevenly distributed; we are only looking at part of the starburst site at UV and optical wavelengths through relatively transparent holes with $E(B-V)$ = 0.8, and the other part is completely obscured at these wavelengths. This explains that $L_{ir}$ is two times greater than the stellar luminosity derived from the broad-band SED analysis.  \\citet{venkatesan06} and \\citet{nozawa07} suggest that dust created in the first SN explosions can be driven through the interior of the SN remnants and accumulated in the SN shells, where second-generation stars may form in compressed cooling gases. Hence, galaxies may be observed as dust-free objects at the very beginning. When much dust formed in SN ejecta, the starburst site would become completely obscured by dust at UV to near-infrared wavelengths. This stage would correspond to extremely dusty objects like DOGs and SMGs. Then, the galactic winds would turned on. \\citet{heckman90} discussed that the galactic winds or superwinds frequently observed in starbursts and ULIRGs will sweep out any diffuse interstellar matter from the starburst site. Such galactic winds will turned on when the thermal energy of the gas heated by SN explosions exceeds the gravitationally binding energy of the gas. According to the models for chemical evolution of elliptical galaxies by \\citet{arimoto87}, galactic winds turn on later in more massive galaxies. The onset of a galactic wind is 350 Myr for a $10^{11} M_{\\sun}$ galaxy, and the metallicity increases up to more than the solar value. These predictions seem to agree well with the age and the metallicity observed in SST J1604+4304, in which the galactic wind would just turn on and dust grains in the opaque shell are pushed and moved outward, creating partially transparent holes in the shell. This picture is consistent with the fact that dusty objects are deemed to be more abundant in massive galaxies than in less massive galaxies, although mid- and far-infrared observations are strongly biased to IR-luminous objects. This observational trend is expected, because galactic winds turn on later and thus the dust shell becomes more opaque for more massive galaxies, resulting in longer time of the dust-obscured stage. This may be the reason why DOGs and SMGs are so infrared-luminous and massive.  \\subsection{Impact by an obscured AGN }  What would happen if the missing 1/2 of the bolometric luminosity comes from an AGN? In this case, the stellar mass ranging 0.1 - 8 $M_{\\sun}$ is reduced to one half of that for the pure star formation. The number of SNe is also reduced by the same amount, while the dust mass does not change. Thus, the dust production rate is increased by a factor of 2, i.e., $0.48 \\pm 0.24 M_{\\sun}$ per SN. This is within a range consistent with the model predictions(\\citealt{nozawa03}; \\citealt{nozawa07}). The metallicity is the same as derived for pure star formation, because the SN ejecta mass is also reduced by a factor of 2. If the AGN shines at the Eddington rate, the mass of the AGN is $2.5 \\times 10^7 M_{\\sun}$."
    },
    "0910/0910.0332_arXiv.txt": {
        "abstract": "Transverse stratification is a common intrinsic feature of astrophysical jets. There is growing evidence that jets in radio galaxies consist of a fast low density outflow at the jet axis, surrounded by a slower, denser, extended jet. The inner and outer jet components then have a different origin and launching mechanism, making their  effective inertia, magnetization, associated energy flux and angular momentum content different as well. Their interface will develop differential rotation, where disruptions may occur. We here investigate the stability of rotating, two-component relativistic outflows typical for jets in radio galaxies. For this purpose, we parametrically explore the long term evolution  of a transverse cross-section of radially stratified jets numerically, extending our previous study where a single, purely hydrodynamic evolution was considered. We include cases with poloidally magnetized jet components, covering hydro and magnetohydrodynamic models. With grid-adaptive relativistic magnetohydrodynamic simulations, augmented with approximate linear stability analysis, we revisit the interaction between the two jet components. We study the influence of dynamically important  poloidal magnetic fields, with varying contributions of the inner component jet to the total kinetic energy flux of the jet, on their non-linear azimuthal stability. We demonstrate that two-component jets with high kinetic energy flux, and an inner jet  effective inertia which is higher than the outer jet effective inertia are subject to the development of a relativistically enhanced, rotation-induced Rayleigh-Taylor type instability. This instability plays a major role in decelerating the inner jet and the overall jet decollimation. This novel deceleration scenario can partly explain the radio source dichotomy, relating it directly to the efficiency of the central engine in launching the inner jet component. The FRII/FRI transition could then occur when the relative kinetic energy flux of the inner to the outer jet grows beyond a certain treshold. ",
        "introduction": "There is strong observational and theoretical evidence that magnetic fields play a crucial role in the acceleration and the collimation of extragalactic jets. Most Active Galactic Nuclei (AGN) jet formation scenarios involve magnetic fields threading a rotating black hole (in the ergosphere) and its accretion disk, thereby removing from them angular momentum, allowing the central black hole to accrete. At least close to the jet launching region, the jet rotation profile persists, reflecting its (general relativistic and/or magneto-rotational) origin. Thus both ingredients, magnetic fields and rotation, are very important in jet formation, as well as in jet propagation and stability.  Moreover, detailed astrophysical jet observations point out that relativistic jets are structured, in a direction perpendicular to the jet axis, typically consisting of a fast spine and slower outer flow. In the case of AGNs, this jet structuring plays an important role in explaining the morphology of the jet high energy radiation \\citep{Ghisellinietal05,Hardcastle06, Jesteretal06,Jesteretal07,Siemiginowskaetal07,Kataokaetal08}, with sometimes clear evidence for a very fast, light inner jet and a heavy slow outer outflow \\citep{Girolettietal04}. Furthermore, observations of the TeV BL Lacertae objects show brightenings and rapid variability in their TeV emission. This variation in their high energy emission implies high Lorentz factor flows occuring at smaller scale, suggesting ultra-relativistic bulk motion of the (inner) jet. At the same time, complementary (radio) observations with VLBI of the pc-scale jet structure indicate a broad, `slowly' (albeit relativistic) moving outflow. In combination, this clearly suggests the presence of a two component jet morphology \\citep{Ghisellinietal05}. Two-component jet structure has also been proposed in more theoretical work, addressing the physics of jet launching, collimation and propagation mechanisms \\citep{BogovalovTsinganos01, Soletal89, Meier03}.  While our jet dynamics computations will be representative for AGN jet conditions, radially structured jet flows are now known to exist in virtually all astrophysical jet outflows. Transversely structured, ultra-relativistic jet-like outflow has been proposed in the context of Gamma Ray Bursts \\citep{Racusinetal08}, to explain the break observed in their afterglow light curve. In the case of stellar outflows, recent observations of some T Tauri jets \\citep{Bacciottietal00, Guntheretal09} also suggest a fast inner outflow bounded by a slow outer outflow. In these young stellar objects, a clear signature of jet rotation around the symmetry axis was detected \\citep{Bacciottietal02, Woitasetal05,Coffeyetal04,Coffeyetal07}, fully supporting scenarios of magneto-centrifugal jet launch and acceleration. Theoretical models of two-component jets in classical T Tauri \\citep{BogovalovTsinganos01,Melianietal06,Cranmer08,Fendt09} then postulate that the inner outflow is turbulent and pressure driven, associated with the young star wind. The inner jet then has a small opening angle, as it is collimated by the outer jet, which is in turn magneto-centrifugally driven from the surrounding disk. The outer disk wind then carries most of the mass loss in the jet. Various authors \\citep{Melianietal06, Fendt09} have demonstrated  using axisymmetric MHD simulations that the outer outflow is self collimated by its intrinsic magnetic field, and that the turbulent inner outflow gets collimated by the outer jet. Furthermore, \\cite{Matsakos08} investigated the topological stability of two-component outflows for young stellar objects, performing extensive numerical simulations to determine whether analytic self-similar models demonstrate robustness in axisymmetric conditions.  Also for relativistic jet simulations, axisymmetry assumptions are often adopted, excluding the development of all non-axisymmetric perturbations. These can address details of how helical field configurations (naturally expected from magneto-centrifugal launch mechanisms) effectively may transport their helicity down the jet beam \\citep{Keppensetal08}, with magnetically aided reacceleration by field compression across internal cross-shocks. While \\cite{Keppensetal08} concentrated on kinetic energy dominated jets, initially Poynting flux dominated jets were simulated by \\cite{Komissarovetal07} in axisymmetric relativistic MHD, finding that the transition to a matter-dominated jet regime occurs very close to the central engine (within $0.01 {\\rm pc}$). Our model computations will therefore assume kinetic energy dominated jets.  As far as the magnetic field topology is concerned, we will restrict ourselves in this paper to purely poloidally magnetized jet components. Our two-component jet model determining our initial conditions can actually allow for helical fields, as explained in Sect.~\\ref{s-model} (we include this more general case here, for future reference in follow-up studies). As indicated before, during the first acceleration phase of AGN jets, magneto-centrifugal mechanisms play an important role and a helical or even strongly toroidal magnetic field is likely produced \\citep{Fendt97,Melianietal06a, McKinney&Blandford09,Komissarovetal07}. \\cite{McKinney&Blandford09} present 3D general relativistic MHD simulations for rapidly rotating black holes producing jets with strong toroidal fields. They find a prominent role of the accreted magnetic field geometry for achieving `stable' jets. Helical or strongly toroidal field topologies can be subject to current-driven kink instabilities \\citep{Begelman98}, with $m=1$ toroidal modes that helically displace the jet axis. This requires full 3D numerical simulations, such as performed by \\cite{BatyKep02} in non-relativistic MHD, or addressed by \\cite{Mizunoetal09} in relativistic MHD for a static force-free equilibrium. Dispersion relations for non-axisymmetric modes and $m=1$ kinks in particular for relativistic MHD were analysed by \\cite{Begelman98} for purely toroidal fields, and electromagnetically dominated, force-free jets were analysed spectrally by \\cite{Istomin&Pariev96} and more recently by~\\cite{Narayanetal09}.  \\begin{figure}[h]  \\begin{center}  {\\resizebox{0.95\\columnwidth}{6cm}{\\includegraphics{Fig1.jpg}}}  \\end{center} \\caption{3D schematic view of the overall AGN disk-jet configuration (indicating the accretion disk and the two-component jet). We model the jet evolution in the transverse plane.}\\label{Fig1} \\end{figure}  In our work, we will restrict attention to 2.5D scenarios with the somewhat unusual assumption of translational invariance along the jet axis. The overall configuration is schematically indicated in Fig.~\\ref{Fig1}, and we simulate a transverse cross-section of the jet at sufficient distance from the two-component jet source, where all three velocity components (axial, azimuthal and radial) are included, but their variation along the jet axis is ignored. Our aim is to investigate all non-axisymmetric instabilities, primarily induced by the (sheared) rotation. This approximation is valid because in the poloidal direction, the flow is supersonic with high Lorentz factor, and then the growth rate of poloidal instabilities is expected to be low. On the other hand, the rotation is subsonic, facilitating the growth of toroidal instabilities. We then address jet stability in a cross section of a rotating two-component jet, initially collimated by thermal pressure and/or poloidal magnetic field. Consecutive snapshots of the cross-sectional evolution can be interpreted as mimicking the jet flow conditions at increasing distance from the source.  Note that this particular assumption allows us to follow both axisymmetric and all non-axisymmetric (including $m=1$) mode development, but does exclude helical mode axis displacement typical for kink modes. Our model therefore mimics jet evolution adequately as long as the radial axis displacement is smaller than the axial wavelength associated with possible $m=1$ kinks. Our assumption does also neglect conical jet expansion, assuming cylindrical propagation. This is justified given the low observed values for jet opening angles.  Since we defer the study of toroidal and helical magnetic field configurations in 2.5D and 3D to later work, we start off with numerically investigating the influence of a purely poloidal magnetic field on the stability of rotating, two-component relativistic jets. Our work complements the studies looking into kink development, by putting the emphasis on the azimuthal variation and on the effect of the two-component jet stratification. As far as a purely poloidal magnetic field topology is concerned, the work by \\cite{Spruitetal97} suggests that jets should be collimated by poloidal magnetic field pressure, rather than by toroidal magnetic field, as toroidal jet magnetic fields can introduce kink instability (but only slow mode growth was found for force-free jets by~\\cite{Narayanetal09}). To justify purely toroidal fields, we can argue that the toroidal magnetic field in the jet acceleration phase can be gradually dissipated involving reconnection, a mechanism which in turn contributes to (axial) jet  acceleration \\citep{Spruit08, Melianietal06}. Indeed, during the acceleration phase a fraction of the Poynting flux (angular momentum) carried by the magnetic field is converted to kinetic energy by internal dissipation/reconnection of the toroidal magnetic field~\\citep{Spruit08, Melianietal06}. In accord with this mechanism, we model the region where the eventual jet rotation is low.     In this paper, we analyse five cases in detail, to determine the effects of  differing poloidal magnetic field configurations in the two-component structure on its long-term stability, and on the overall deceleration of the jet as it propagates away from the central source regions. The role of poloidal magnetic field is in these cases most prominent in its added effect on total pressure and effective fluid inertia, and this is shown to play a prominent role in the two-component jet stability. ",
        "conclusions": "We examined five configurations of magnetized two-component jets. All share the same density ratio between inner and outer component and identical rotation profiles. The magnetic and thermal pressure configuration in each model differs, though. This in turn translates to different distributions of the (total, fixed) kinetic energy flux over the inner and outer jet component. In fact, in case (A) the kinetic energy flux in the inner component is about $10\\%$, while it is $0.7\\%$ for case (B1) and $0.5\\%$ in case (B2), and around $38\\%$ for cases (C) and (D). For the outer component, we thus have $90\\%$ in case (A), $99.3\\%$ in case (B1) and $99.5\\%$ in case (B2), and only $62\\%$ for cases (C) and (D). The most important difference between the two model categories is then: cases (A, C, D) have an inner jet component with higher inertia $\\gamma^2\\,\\rho\\,h+B_{\\rm z}^2$ than their outer jet component, while cases (B1) and (B2) have an inner jet component inertia which ends up lower than in the outer jet component. From the detailed analysis of the simulations, as well as from the approximate stability analysis, this criterion distinguishes between cases where relativistically enhanced Rayleigh-Taylor modes ultimately dominate the evolution, leading to complete mixing of both components and inner jet deceleration. This is quantified most clearly by showing the time evolution of the mean Lorentz factor over the inner jet region for all cases in Fig.~\\ref{Lorentzfactor_evolution}. This requires a clear criterion to distinguish inner versus outer jets in the turbulent evolutions. In cases (A), (C) and (D), we locate the outer jet component as having a Lorentz factor $2.5<\\gamma<3.5$ and effective polytropic index $\\Gamma_{\\rm eff}>3/2$. The inner component is the region defined by a Lorentz factor $\\gamma\\ge 3.5$ and effective polytropic index $\\Gamma_{\\rm eff}\\le 3/2$. In the cases (B1) and (B2), the effective polytropic index in the inner and outer jet can be locally of the same order there, hence we use that the inner jet is magnetized. In the cases (B1) and (B2), the inner component jet and shear region are not totally mixed during the evolution, since the inner component is compressed and has higher magnetization and Lorentz factor. Thus to distinguish the inner component jet we put a condition on magnetic field strength $B_{\\rm z}>B_{\\rm z, initial}/2$, where $B_{\\rm z, initial}$ is the magnetic strength assumed initially. Under these precise quantifications of inner/outer jet regions, Fig.~\\ref{Lorentzfactor_evolution} demonstrates clearly how stable cases (for the relativistically enhanced Rayleigh-Taylor modes) remain at high speed, while unstable cases decelerate. Using the same means to distinguish inner versus outer jet regions at all times, we can quantify the inner jet radius for all cases, as well as the total jet radius for all cases. These are shown in Figs.~\\ref{Inradius_evolution}-\\ref{Radius_evolution}, and quantify the decollimation effects discussed in Section~\\ref{Simulations}.   During the entire evolutions, the toroidal and radial speeds remain weak as we have typically a maximal $V_{R}<0.01$ and $V_{\\varphi}<0.04$. This means that the contribution of the laboratory frame charge separation force $\\rho_{e}{\\vec{E}}$ to the Lorentz force is negligible at all times. Under these conditions, in both radial and toroidal directions, the contribution of magnetic energy to fluid inertia is $1/\\gamma^2$ weaker than the contribution of the magnetic pressure to the total pressure. This explains why cases (B1) and (B2) are then more stable than the other cases. Despite the fact that case (B2) has similar axial total pressure than  other cases (A, C, D), the low contribution of thermal energy to total pressure makes the effective inertia ratio between the inner and outer jet low. This two-component jet is then stable against the relativistically enhanced Rayleigh-Taylor instability. Extraction of angular momentum and energy from the inner jet to the shear shell is less efficient in cases (B1) and (B2) than in all other cases. As a result, in cases (B1) and (B2) the inner relativistic jet with high Lorentz factor $\\gamma_{\\rm in}\\sim 20$  persists. In the other cases, the kinetic energy flux in the inner jet was initially relatively high, making them unstable and leading to deceleration to Lorentz factors around $8$.   We investigated the stability of two-component jets  beyond the launching region, where both components are rotating differently and a clear two-component structure exists. We initialized this model in accord with magneto-centrifugal models for jet generation, also using the observed analogy between radio source jets and two-component jets in young stellar objects, where the rotation within the jet can actually be observed. We performed five very high resolution simulations of magnetized two-component jets with various magnetization and kinetic energy flux stratifications. The two-component jets with a low inner kinetic energy flux contribution are more stable and remain relativistic for long distances, whereas jets with a highly contributing inner jet to the total jet kinetic energy flux, are subject to a relativistic Rayleigh-Taylor type instability. This instability turns out to be very efficient to decollimate and decelerate the inner jet. Jets that are subject to this instability become turbulent after propagating for a distance of about $30 {\\rm pc}$.  This new result on two-component jet models is important because it can explain the classification of radio sources in Fanaroff-Riley I/II categories according the energy stratification of the inner jet. This ultimately relates to the jet launch region and the properties of the inner accretion disk. In fact, by analogy between the FRI/FRII classification  and the results of our model, an FRI jet would correspond to a two-component jet with a high energy flux contribution from the inner jet, whereas the FRII jet corresponds to relatively low energy fluxes in the inner jet. The model we propose here to explain the FRI/FRII dichotomy is different from the model we proposed earlier~\\citep{Melianietal08} where the transition occurs due to external density stratification. That model explains the group of peculiar ``HYbrid MOrphology Radio Source\" (HYMORS) \\citep{GopalKrishna&Wiita00} which appear to have an FR II type on one side and an FR I type diffuse radio lobe on the other side of the active nucleus. Since the launch conditions on each side are presumably similar in these kind of radio sources, the different Fanaroff-Riley morphologies on either side must be attributed to the change in the properties of the ambient media, as shown convincingly in~\\cite{Melianietal08}. The results of the present paper nicely complement these earlier findings with a quantifiable role of the central engine contribution.  We currently continue this study of the interaction between two component jets in full 3D. We thereby intend to explore the relative influence of azimuthal versus longitudinal instabilities for realistic multi-component jets. Another extension is to allow for aximuthal magnetic fields, in accord with the initial profiles as given generally in this paper.  It then rremains to be shown that (1) the newly discovered instability persists in 3D hydro and magnetohydrodynamic configurations, where the potential role of axial mode development is incorporated, and introduces helical jet axis displacements; (2) how the instability gets modified (stabilized or destabilized) by the inclusion of toroidal field components, first in 2.5D neglecting helical axis displacements, and consecutively in 3D, where also current-driven kinks may occur."
    },
    "0910/0910.5721_arXiv.txt": {
        "abstract": "One of the primary means of determining the spin $a$ of an astrophysical black hole is by actually measuring the inner radius $r_\\mathrm{in}$ of a surrounding accretion disk and using that to infer $a$. By comparing a number of different estimates of $r_\\mathrm{in}$ from simulations of tilted accretion disks with differing black-hole spins, we show that such a procedure can give quite wrong answers. Over the range $0 \\le a/M \\le 0.9$, we find that, for moderately thick disks ($H/r \\sim 0.2$) with modest tilt ($15^\\circ$), $r_\\mathrm{in}$ is nearly independent of spin. This result is likely dependent on tilt, such that for larger tilts, it may even be that $r_\\mathrm{in}$ would increase with increasing spin. In the opposite limit, we confirm through numerical simulations of untilted disks that, in the limit of zero tilt, $r_\\mathrm{in}$ recovers approximately the expected dependence on $a$. ",
        "introduction": "\\label{sec:intro} Classical, astrophysical black holes have only two defining properties: mass and angular momentum (or spin). As in other astrophysical contexts, the mass of a black hole can straightforwardly be determined by the application of Kepler's Third Law, provided the black hole is orbited by another observable object and the distance to the system is reasonably well known (at least to within a factor of $\\sin i$, where $i$ is the inclination of the binary orbit relative to the observer's line-of-sight).  The spin $a=Jc/GM$, on the other hand, is a very different matter. So far, no direct measure of spin has been proposed, leaving only indirect means of inferring it. A commonly used method is based on making some measure of the ``inner radius\" of the accretion disk. In general relativity, stable circular orbits are not permitted all the way down to the ``surface'' of the black hole, but are instead restricted to lie outside the marginally stable limit $r_\\mathrm{ms}$. Beginning with \\citet{nov73}, it has commonly been assumed that the accretion disk will observe a similar restriction and truncate at $r_\\mathrm{ms}$. Because $r_\\mathrm{ms}$ has a well-known monotonic dependence on the spin of the black hole \\citep{bar72}, its association with an observed inner radius of the accretion disk is then used to infer the spin of the black hole. The observation of the inner radius can come either from modeling of the continuum emission in the thermally dominant state \\citep[e.g.][]{sha06,dav06} or from measuring relativistically broadened reflection features \\citep[e.g.][]{wil01}.  However, it has been shown that there are many difficulties with such procedures. First, one has many choices for how to define the effective inner radius $r_\\mathrm{in}$ of the accretion disk, and some of these can vary considerably (factors of 2-3) from $r_\\mathrm{ms}$ \\citep{kro02}. Even if one restricts oneself to such observationally relevant measures as the ``radiation edge'' (the innermost radius from which significant luminosity emerges) or ``reflection edge'' (the smallest radius capable of producing significant X-ray reflection features), significant deviations may be observed \\citep{kro02, bec08}. For instance, \\citet{ago00} showed that significant emission can emerge from inside $r_\\mathrm{ms}$, thus precluding a direct correlation with $r_\\mathrm{in}$. Similarly \\citet{rey97} showed that the reflection edge is not necessarily tied to $r_\\mathrm{ms}$.  Nevertheless, a strong result for untilted disks (disks that lie approximately in the symmetry plane of their black hole spacetimes and have their angular momenta aligned with the spins of their black holes) is that $r_\\mathrm{in}$ should gradually decrease with increasing black hole spin, i.e. as $a/M\\rightarrow 1$. Therefore, it might at least be hoped that a study of $r_\\mathrm{in}$ for a large ensemble of accreting black holes might yield some information about the range and distribution of black hole spins represented in the ensemble.  However, we show in this Letter that there is a further complicating factor. Based on numerical simulations, we show that the effective inner radius of tilted disks (disks that do not have their angular momenta aligned with the spins of their black holes) yield very different results for $r_\\mathrm{in}$ than their untilted counterparts for all dynamical measures we have considered. A disk could be tilted for many reasons. In stellar mass binaries, the orientation of the outer disk is fixed by the binary orbit, whereas the orientation of the black hole is determined by how it became part of the system. If the black hole formed from a member of a preexisting binary through a supernova, then the black hole could be tilted if the explosion were asymmetric. If the black hole joined the binary through multi-body interactions, such as binary capture or replacement, then there would have been no preexisting symmetry, so the resulting system would nearly always harbor a tilted black hole. This same argument can be extended to AGN in which merger events reorient the central black hole or its fuel supply and result in repeated tilted configurations. Tilted disks are subject to additional torques due to differential Lense-Thirring precession \\citep{bar75}.  Our results indicate that for moderately thick ($H/r \\sim 0.2$), tilted ($15^\\circ$) disks, $r_\\mathrm{in}$ is independent of spin, or, in some measures, actually {\\em increases} with increasing black hole spin. These results could have important consequences for observational efforts to constrain black hole spins. ",
        "conclusions": "\\label{sec:discussion} We have explored four possible measures of the effective inner radius of simulated black-hole accretion disks, based on surface density, inflow and turbulent velocities, and magnetic field structure. With all four measures, we found two consistent trends: 1) for untilted simulations $r_\\mathrm{in}$ closely follows the analytic form of $r_\\mathrm{ms}$; whereas 2) for $15^\\circ$ tilted simulations $r_\\mathrm{in}$ shows a flat or even increasing trend as $a/M\\rightarrow 1$.  The fact that Figures \\ref{fig:density} - \\ref{fig:alpha} all show remarkable similarities suggests that our results are not due to a poor choice of definition for $r_\\mathrm{in}$. Also, although we have ignored the dependence of $r_\\mathrm{ms}$ on inclination in making our plots, Figure 5 of \\citet{fra07} shows that the deviation would be quite small for $i = 15^\\circ$, as would be appropriate for this work.  We can get some estimate of the uncertainties associated with our results by comparing data obtained from the two 515H simulations, one done on the spherical-polar grid and one on the cubed-sphere. For the first three measures of $r_\\mathrm{in}$, where we were able to get estimates using both simulations, they agree remarkably well. Furthermore, for most of the measures the trend lines fit the data well, suggesting small uncertainties.  We conclude that the discrepancy of $r_\\mathrm{in}$ between tilted and untilted simulations is a robust result, and must, therefore, be rooted in some physical process. The most likely agent of such a process would seem to be the standing shocks aligned along the line of nodes between the black hole symmetry plane and disk midplane \\citep{fra08}. These features are not present in simulations of untilted disks. Such non-axisymmetric shocks can be quite efficient at extracting angular momentum from the gas flow, which would reduce the effect of the black hole spin on $r_\\mathrm{in}$.  In this work we have only considered two values of initial tilt: $\\beta_i = 0$ and $15^\\circ$. It stands to reason, however, that as $\\beta_i \\rightarrow 0$ the inner radius would approach the untilted values obtained in this work. On the other hand, for $\\beta_i > 15^\\circ$, it may well be that the discrepancies in $r_\\mathrm{in}$ would be even larger. In such a case, $r_\\mathrm{in}$ might truly increase with increasing spin.  The results in this Letter then imply that, at least for moderately thick accretion disks ($H/r \\gtrsim 0.1$), measurements of $r_\\mathrm{in}$ are {\\em not} reliable predictors of $a$, unless one can independently confirm that the disk is not tilted. This could pose a problem for recent attempts to determine spin using so-called ``Hard'' state observations \\citep{mil06, rei08}, where the flow is generally taken to be thick.  For very thin disks ($H/r \\lesssim 0.01$), the effect of tilt is expected to be much different. Such disks are expected to be subject to the Bardeen-Petterson effect \\citep{bar75}, where the midplane of the disk at small radii would be aligned with the symmetry plane of the black hole through the competing action of Lense-Thirring precession and viscous diffusion, while possibly leaving an outer disk that is tilted at large radii. Since the inner disk would be aligned with the symmetry plane of the black hole, the effective inner radius should be the same as for an untilted disk. This would apply for disks in the ``Soft'' or ``Thermally-Dominant'' state, with luminosities well below the Eddington limit, which are likely to be thin. Attempts to estimate the black hole spin by modeling the disk in this state may be unaffected by the cautions raised in this Letter.  However, for sources with disk luminosities near or above the Eddington limit \\citep[e.g.][]{mid06}, the disk should again be thick, and the effect discussed here will apply.  Even if the disk is thin, we reiterate that it is still important to independently fit for the inclination of the inner disk, since this may not be the same as the binary inclination.  Other methods of determining black-hole spin, such as using quasi-periodic oscillation (QPO) frequencies, may not be affected by our results. For instance, QPO models based on resonant frequencies attributable to the black hole spacetime itself \\citep[e.g.][]{abr01} would likely not be strongly affected. In fact, the weak dependence of $r_\\mathrm{in}$ on $a$ detailed in this Letter actually helps one model of low-frequency QPOs based on the precession of a radially extended thick disk, by ensuring that the maximum frequency would not exceed $\\sim 10$ Hz regardless of black hole spin, as observed \\citep{ing09}. However, to date, no consensus model for QPOs has emerged, leaving us with few options for unambiguously determining black-hole spin.  To end, we emphasize that it would be wrong to conclude from this paper that all tilted black hole accretion disks should appear to have $r_\\mathrm{in} \\approx 6 M$. Remember, the normalizations used in each of our plots were chosen simply to give values for the untilted disks that were reasonably similar to the values for $r_\\mathrm{ms}$. There is no other physical motivation for choosing those normalizations, and a different choice would have led to different numerical values for $r_\\mathrm{in}$. Further, we can not really say how the dynamically determined $r_\\mathrm{in}$ in this work would compare to values of $r_\\mathrm{in}$ measured from continuum modeling or iron-line fitting without doing more work. Specifically, we would need to generate synthetic disk images from the simulation data, and measure a radiation or reflection edge. Such work is already underway using the relativistic ray tracing code of \\citet{dex09}. Even then, there is the problem that the original simulations did not properly account for all of the important physical processes occurring in the disk, notably heating due to dissipation of turbulent energy and radiative cooling. That will have to await future simulations.  What {\\em can} be concluded from this paper is that all tilted disks (with a tilt around $15^\\circ$) should appear to have the {\\em same} $r_\\mathrm{in}$, independent of spin. Our results may further imply that measurement of a small value for $r_\\mathrm{in}$ (without defining small) would require a rapidly rotating black hole {\\em and} an untilted disk, whereas a large value of $r_\\mathrm{in}$ would {\\em not} necessarily mean that the black hole is spinning slowly; it could simply be tilted."
    },
    "0910/0910.2337_arXiv.txt": {
        "abstract": "Accretion of dark energy onto black holes will take place when dark energy is not a cosmological constant. It has been proposed that the time evolution of the mass of the black holes in binary systems due to dark energy accretion could be detectable by gravitational radiation. This would make it possible to use observations of black hole binaries to measure local dark energy properties, e.g., to determine the sign of $1+w$ where $w$ is the dark energy equation of state. In this Letter we show that such measurements are unfeasible due to the low accretion rates. ",
        "introduction": "Experimental evidence shows that the universe currently undergoes an accelerated expansion driven by an unknown form of energy, dubbed dark energy. This dominant energy component can -- at least phenomenologically -- be described as a perfect fluid with density $\\rho_{de}$ and a (possibly time dependent) equation of state $w\\equiv p_{de}/\\left(\\rho_{de}c^{2}\\right)$, where $p_{de}$ is the dark energy pressure. Current measurements are consistent with the dark energy density being approximately 70\\,\\% of the critical density of the universe, $\\rho_{de}\\sim 10^{-29}\\,\\mbox{g cm}^{-3}$, and the equation of state being close to $w=-1$, i.e., the value corresponding to a cosmological constant. Since the density is very low and we do not expect dark energy to cluster to any large extent, all the information we have on dark energy is through its impact on the largest scales of the universe via the expansion rate of the universe. One of the foremost experimental and theoretical tasks in cosmology is to pin down the equation of state $w$ of the dark energy, specifically to constrain or detect any deviations from the cosmological constant value of $w=-1$.  It has been proposed that dark energy can be detected also on smaller scales \\cite{mersini,wang}. Any such local measurements would constitute a powerful confirmation of the existence of dark energy. In \\citet{mersini}, the possibility to measure local properties of dark energy using gravitational radiation from binary systems of supermassive black holes is investigated. The basic idea is that accretion of dark energy by black holes will lead to a mass change when dark energy is not a cosmological constant. The mass accretion would affect the dynamics of the system and in turn alter the waveform of the gravitational radiation produced by the binaries during the inspiraling phase. Measuring this waveform, through LIGO or LISA, would then make it possible to constrain the values of $w$ and $\\rho_{de}$. Alternatively, the change in orbital period of the system could be measured using electromagnetic radiation, if present. In this Letter we claim that this effect is negligible and that this form of local measurement of $\\rho_{de}$ and $w$ is not possible.  In Section 2 of the Letter, we compare the mass accretion due to the dark energy with that of the interstellar medium and dark matter. In Section 3 we look at the radial separation needed for dark energy accretion to dominate over the effect over gravitational radiation in a binary system of two supermassive black holes. In both cases we show that dark energy accretion cannot be considered a measurable effect, even for a binary system. In Section 4 we consider two concrete examples to validate this conclusion. ",
        "conclusions": "We have investigated the effect of dark energy accretion in binary supermassive black hole systems. When comparing the mass change due to accretion from the interstellar medium, dark matter and dark energy, we find that the effect of dark energy accretion for all practical purposes is negligible. We also compare the effects from dark energy accretion and the loss of energy through gravitational wave emission. At the separations between the binary constituents needed for dark energy accretion to dominate over gravitational radiation, and under the assumption that the rotation of dark energy in the vicinity of the binary system is not greatly exceeding that of the binary system, we find the radial change induced by the accretion to be too small to be observable. Thus, even in an idealized setting with no gas accretion, the effect of dark energy accretion during the inspiraling phase does not have a measurable impact on the dynamics of the system. Doing local measurements of the equation of state of dark energy, as proposed in \\citet{mersini}, can therefore unfortunately not be considered possible."
    },
    "0910/0910.5818_arXiv.txt": {
        "abstract": "The radioscience experiment is one of the on board experiment of the Mercury ESA mission BepiColombo that will be launched in 2014. The goals of the experiment are to determine the gravity field of Mercury and its rotation state, to determine the orbit of Mercury, to constrain the possible theories of gravitation (for example by determining the post-Newtonian (PN) parameters), to provide the spacecraft position for geodesy experiments and to contribute to planetary ephemerides improvement. This is possible thanks to a new technology which allows to reach great accuracies in the observables range and range rate; it is well known that a similar level of accuracy requires studying a suitable model taking into account numerous relativistic effects. In this paper we deal with the modelling of the space-time coordinate transformations needed for the light-time computations and the numerical methods adopted to avoid rounding-off errors in such computations. ",
        "introduction": "\\label{intro}  BepiColombo is an European Space Agency mission to be launched in 2014, with the goal of an in-depth exploration of the planet Mercury; it has been identified as one of the most challenging long-term planetary projects. Only two NASA missions had Mercury as target in the past, the Mariner 10, which flew by three times in 1974-5 and Messenger, which carried out its flybys on January and October 2008, September 2009 before it starts its year-long orbiter phase in March 2011.  The BepiColombo mission is composed by two spacecraft to be put in orbit around Mercury. The radioscience experiment is one of the on board experiments, which would coordinate a gravimetry, a rotation and a relativity experiment, using a very accurate range and range rate tracking. These measurements will be performed by a full 5-way link (\\cite{iess01}) to the Mercury orbiter; by exploiting the frequency dependence of the refraction index, the differences between the Doppler measurements (done in Ka and X band) and the delay give information on the plasma content along the radiowave path. In this way most of the measurements errors introduced can be removed, improving of about two orders of magnitude with respect to the past technologies. The accuracies that can be achieved are $10$ cm in range and $3 \\times 10^{-4}$ cm/s in range rate.  How we compute these observables?  For example, a first approximation of the range could be given by the formula \\begin{equation} r=|{\\bf r}|=|({\\bf x}_{\\rm sat}+{\\bf x}_{\\rm M})-({\\bf x}_{\\rm EM}+ {\\bf x}_{\\rm E}+{\\bf x}_{\\rm ant})| \\,\\, , \\label{eq:new} \\end{equation} which models a very simple geometrical situation (see Figure~\\ref{fig:range}). The vector ${\\bf x}_{\\rm sat}$ is the mercurycentric position of the orbiter, the vector ${\\bf x}_{\\rm M}$ is the position of the center of mass of Mercury (M) in a reference system with origin at the Solar System Barycenter (SSB), the vector ${\\bf x}_{\\rm EM}$ is the position of the Earth-Moon center of mass in the same reference system, ${\\bf x}_{\\rm E}$ is the vector from the Earth-Moon Barycenter (EMB) to the center of mass of the Earth (E), the vector ${\\bf x}_{\\rm ant}$ is the position of the reference point of the ground antenna with respect to the center of mass of the Earth.  \\begin{figure}[h] \\begin{center} \\includegraphics[width=6cm]{fig_range.eps} \\end{center} \\caption{\\footnotesize Geometric sketch of the vectors involved in the computation of the range. SSB is the Solar System Barycenter, M is the center of Mercury, EMB is the Earth-Moon Barycenter, E is the center of the Earth.} \\label{fig:range} \\end{figure}   Using (\\ref{eq:new}) means to model the space as flat arena ($r$ is an Euclidean distance) and the time as absolute parameter.  This is obviously not possible because it is clear that, beyond some threshold of accuracy, these quantities have to be formulated within the framework of Einstein's theory of gravity (general relativity theory, GRT). Moreover we have to take into account the different times at which the events have to be computed: the transmission of the signal at the transmit time ($t_t$), the signal at the Mercury orbiter at the time of bounce ($t_b$) and the reception of the signal at the receive time ($t_r$).  Formula (\\ref{eq:new}) could be a good starting point to construct a correct relativistic formulation; with the word ``correct'' we do not mean all the possible relativistic effects, but the effects that can be measured by the experiment. This paper deals with the corrections to apply to this formula to obtain a consistent relativistic model for the computations of the observables and the practical implementation of such computations.  In Section \\ref{sec:refsys} we discuss the relativistic four-dimensional reference systems used and the transformations adopted to make the sums in (\\ref{eq:new}) consistent; according to \\cite{soffel03}, with ``reference system'' we mean a purely mathematical construction, while a ``reference frame'' is a some physical realization of a reference system. The relativistic contribution to the time delay due to the Sun's gravitational field, the Shapiro effect, is described in Section \\ref{sec:shap}.  Section \\ref{sec:lt} deals with the theoretical procedure to compute the light-time (range) and the Doppler shift (range rate).  In Section~\\ref{sec:num} we discuss the practical implementation of the algorithms showing how we eliminate rounding-off problems.  The equations of motion for the planets Mercury and Earth, including all the relativistic effects (and potential violations of GRT) required to the accuracy of the BepiColombo radioscience experiment have already been discussed in \\cite{mil09b}, thus this paper concentrates on the computation of the observables. ",
        "conclusions": "\\label{sec:conc}  By combining the results of the previous paper (\\cite{mil09b}), and of this one, we have completed the task of showing that it is possible to build a consistent relativistic model of the dynamics and of the observations for a Mercury orbiter tracked from the Earth, at a level of accuracy and self-consistency compatible with the very demanding requirements of the BepiColombo radioscience experiment.  In particular, in this paper we have given the algorithm definitions for the computation of the observables range and range rate, including the reference system effects and the Shapiro effect. We have shown which computation can be performed explicitly and which ones need to be obtained from an iterative procedure. We have also shown how to push these computations, when implemented in a realistic computer with rounding-off, to the needed accuracy level, even without the cumbersome usage of quadruple precision. The list of ``relativistic corrections'', assuming we can distinguish their effects separately, is long, and we have shown that many subtle effects are relevant to the required accuracy. However, in the end what is required is just to be fully consistent with a post-Newtonian formulation to some order, to be adjusted when necessary. Interestingly, the high accuracy of BepiColombo radio system may require implementation of the second post-Newtonian effects in range."
    },
    "0910/0910.5735_arXiv.txt": {
        "abstract": "Current and future weak lensing surveys will rely on photometrically estimated redshifts of very large numbers of galaxies. In this paper, we address several different aspects of the demanding photo-z performance that will be required for future experiments, such as the proposed ESA Euclid mission.  It is first shown that the proposed all-sky near-infrared photometry from Euclid, in combination with anticipated ground-based photometry (e.g. PanStarrs-2 or DES) should yield the required precision in individual photo-z of $ \\sigma_{z} (z) \\leq 0.05(1+z)$ at $ I_{AB} \\leq 24.5$.  Simple a priori rejection schemes based on the photometry alone can be tuned to recognise objects with wildly discrepant photo-z and to reduce the outlier fraction to $ \\leq 0.25\\%$ with only modest loss of otherwise usable objects.  Turning to the more challenging problem of determining the mean redshift $ \\langle z\\rangle$ of a set of galaxies to a precision of $ |\\Delta_{\\langle z \\rangle}|\\leq 0.002(1+z)$ we argue that, for many different reasons, this is best accomplished by relying on the photo-z themselves rather than on the direct measurement of $\\langle z\\rangle$ from spectroscopic redshifts of a representative subset of the galaxies, as has usually been envisaged. We present in an Appendix an analysis of the substantial difficulties in the latter approach that arise from the presence of large scale structure in spectroscopic survey fields.  A simple adaptive scheme based on the statistical properties of the photo-z likelihood functions is shown to meet this stringent systematic requirement.  We also examine the effect of an imprecise correction for Galactic extinction on the photo-z and the precision with which the Galactic extinction can be determined from the photometric data itself, for galaxies with or without spectroscopic redshifts. We also explore the effects of contamination by fainter over-lapping objects in photo-z determination.  The overall conclusion of this work is that the acquisition of photometrically estimated redshifts with the precision required for Euclid, or other similar experiments, will be challenging but possible. ",
        "introduction": "Large scale mapping of the weak lensing shear field in three dimensions is emerging as a potentially very powerful cosmological probe  \\citep{Peacock2006, DETF}. Weak lensing has the advantage of directly tracing the mass distribution, thereby bypassing much of the complex astrophysics of the baryon component that underpin most of the other probes and which may well dominate the systematic uncertainties in them. In contrast, the underlying physics of weak lensing is extremely simple, and the challenges are primarily on the observational side, particularly the accurate measurement of the weak lensing distortion and the estimation of distances to very large numbers of faint galaxies.  A weak lensing survey of half the sky $ \\rm (20,000\\, deg^{2})$ to a depth of $ I_{AB} \\sim 24.5$ and with a PSF of $ \\sim$ 0.2 arcsec, forms a major part of the proposed ESA Euclid mission\\footnote{http://www.ias.u-psud.fr/imEuclid \\\\ http://sci.esa.int/science-e/www/area/index.cfm?fareaid=102}. Euclid had its origins in two proposals submitted for the first round of the ESA Cosmic Visions 2015-2025 competition, the DUNE imaging survey \\citep{DUNE2006} and the SPACE spectroscopic survey \\citep{SPACE2009}.  The combination of the two surveys, plus the anticipated improved information on the Cosmic Microwave Background from Planck\\footnote{www.rssd.esa.int/Planck} offers dramatic improvements in our knowledge of the entire dark sector, including the definition of the dark matter power spectrum, the dark energy equation of state parameter $ w$, as well as much else.  Application of weak lensing for cosmology requires at least a statistical knowledge of the distances, i.e. redshifts, of large numbers of individual galaxies. At $ I_{AB} <  24.5$, there are about 2.5 billion galaxies in the Euclid $ 2\\pi$ sr survey area and so, realistically, reliance must be made on photometrically estimated redshifts (hereafter \\textit{photo-z}).  \\subsection{Required redshift precision for precision cosmology with weak lensing}  Several papers have discussed the redshift precision that is needed for weak lensing analyses to enable the full potential of this approach to be exploited \\citep{Amara&Refrigier2007, Ma&HUetal2006, Abdalla2008}.  In the lensing tomography approach \\citep{Hu_Tomography}, individual galaxies are binned into a number of redshift bins. The shear signal is extracted from the cross-correlation of the shape measurements of individual galaxies in different redshift bins.  These correlated alignments then give information (with some distance weighting function) on the mass distribution between the observer and the nearer of the two redshift bins.  Redshift information for the galaxies is required at two conceptually distinct steps: first, the construction of the redshift bins used for the cross-correlation analysis to extract the weak lensing signal and, second, the estimation of the mean redshift of the galaxies in a given bin, which is required to map the results onto cosmological distance and thereby extract the cosmological parameters. It is, of course, possible to do a similar correlation analysis with unbinned data \\citep{2005PhRvD..72b3516C, kitching2008}, but for our purposes this distinction is unimportant.  The cross-correlation between different redshift bins is undertaken to exclude any galaxy pairs that may be physically associated, i.e. have the same distance.  This is to avoid the possibility that physical processes operating around individual galaxies may produce an intrinsic alignment of the galaxies that may be mistaken for the coherent alignment produced by the weak lensing of the foreground mass distribution. The required accuracy of the individual photo-z for the bin-construction task is set by the need to exclude overlaps in the N(z) of individual bins (or the probability distribution for individual galaxies) and thereby remove physically close pairs \\citep{2002A&A...396..411K}. This typically sets a requirement on the precision of individual photo-z of about $ \\sigma_{z} = 0.05(1+z)$ \\citep{bridle&King2007}.  There is a second type of intrinsic alignment effect \\citep{2004PhRvD..70f3526H}, whereby the shape of the further of a given galaxy pair may be affected, through lensing effects, by the shape of the matter distribution around the nearer galaxy, which is likely to be correlated with the visible shape of that galaxy, thereby again producing a correlated alignment of the two galaxies that is unfortunately nothing to do with the lensing signal from the common foreground. \\cite{Jochimi&Schneider2008},~\\cite{Joachimi&Schnider2009} have shown that it is possible to implement a nulling approach to eliminate this second intrinsic alignment signal, which again requires a priori knowledge of individual redshifts.  Once the weak lensing signal is extracted, accurate knowledge of the redshift of the galaxies, as with any cosmological probe gives, amongst other parameters, information on the angular diameter distance $ D_{\\theta}(z)$. The sensitivity of $ D_{\\theta}(z)$ to the relevant cosmological parameters ($ \\Omega_{m}$, $ \\Omega_{\\Lambda}$, $w$ etc) gives the required accuracy in the mean redshifts that are required to achieve a given precision on the parameters. As an example, \\cite{Peacock2006} have shown that a precision of 1$ \\%$ in $w$ requires a typical precision in the mean redshift of about 0.2$ \\%$ in  $\\langle z\\rangle$. The Euclid goal is a 2$ \\%$ precision in $w$ (independent of priors), requiring a precision of order 0.002(1+z) in $\\langle z\\rangle$. This simple approach is confirmed by extensive analysis of the Fisher matrices \\citep{Amara&Refrigier2007, Ma&HUetal2006}. It is generally the case that if the mean redshift of a bin is defined accurately enough, then the higher moments of N(z) within the bin will also have been sufficiently determined. Of course, systematic biases in $\\langle z\\rangle$  that vary smoothly with redshift are particularly troublesome as they will mimic the effect of changing the cosmological parameters.  In summary, in order to reach the Euclid performance, we require a statistical (random) r.m.s. precision of order $0.05(1+z)$ per galaxy (for the correlation analysis), and a systematic precision in the mean z in each bin of order 0.002(1+z). These are both quite demanding requirements, and together with the shape measurement itself ($ \\delta \\gamma \\sim 3 \\times 10^{-4}$ - \\citep{GREAT08, 2008MNRAS.391..228A}), they represent one of the observational challenges that lie along the path to enabling precision cosmology with weak lensing.  Fortunately, there are some mitigating features of weak lensing analysis. For instance, the analysis is robust (aside from root-n statistics) to the exclusion of individual galaxies, provided only that the exclusion is unrelated to their shapes.  One is free therefore to reject galaxies that are likely to have poor photo-z provided that they can be recognized a priori, i.e. from the photometric data alone.  \\subsection{Challenges for the spectroscopic calibration of N(z)}  Given the stringent requirements on the systematic error in the mean redshift $\\langle z\\rangle$ of a particular bin, one approach \\citep{Abdalla2008} is to define the $N(z)$ and mean $\\langle z\\rangle$ through the acquisition of spectroscopic redshifts for a representative subset of the galaxies.  This direct sampling approach is certainly the most conservative, but will be very challenging, in practice, for the following reasons.  First, one clearly requires very large numbers of spectroscopic redshifts.  If we have a total redshift interval of $ \\Delta z$, split into m bins, then the number of spectroscopic redshifts N will be of order:  \\begin{equation} N\\, \\sim  \\, m^{-1} (\\Delta z/ \\sigma_{\\langle z\\rangle})^{2} \\end{equation} This assumes that the photo-z are perfect, and that there are no outliers.  This leads trivially \\citep{Amara&Refrigier2007} to a requirement for $ 10^{5}-10^{6}$ spectroscopic redshifts.  Secondly, these spectroscopic redshifts must be fully representative of the underlying sample.  Any biases in the sampling of the bin or, even harder to reliably quantify, the almost inevitable biases in the ability to secure a reliable spectroscopic redshift, must be dealt with via a possibly complex and inevitably somewhat uncertain weighting scheme \\citep{limaetal2008}. It should be noted that current routine spectroscopic surveys of typical faint galaxies do not even approach 100$ \\%$ completeness, even at brighter levels.  One of the best to date is the zCOSMOS survey \\citep{Lilly2007} on relatively bright $ I_{AB} < 22.5$ galaxies which yields, at present, a 99$ \\%$ secure redshift for 95$ \\%$ of galaxies at its optimum $ 0.5 < z < 0.8$ \\citep{Lilly2009}. Most other surveys are significantly less complete.  Even more invidious are the effects of large scale structure in the spectroscopic survey fields, often called cosmic variance. Our own semi-empirical analysis (see the Appendix) of the COSMOS mock catalogues \\citep{kitzbichler&White2007} shows that, in a given patch of sky, the N(z) at $ I_{AB} \\sim 24$ becomes dominated by cosmic variance as soon as a rather small number of galaxies have been observed.  The precise number depends on the field of view of the spectrograph, but is typically about 20-100 for spectroscopic fields of order 0.02-1 square degrees, i.e. a sampling rate of only a few percent.  This means that the spectroscopic survey must be split up over a very large number of independent fields and that to get $ 10^{6}$ redshifts that are Poisson variance dominated one must effectively cover the whole sky in a sparse sampled way.  This is unlikely to be efficient with the large telescopes needed for such faint object spectroscopy. A similar concern comes from the effects of Galactic extinction and reddening, which are likely, even when corrected for, to require spectroscopic sampling across the full range of Galactic [b,l].  These difficulties prompt consideration of other approaches, and in particular, that of placing greater reliance on the photo-z themselves, not only to construct the bins, but also to define their $\\langle z\\rangle$ with small systematic error.  \\subsection{Using photo-z for construction of N(z)}  The performance of photometric redshifts is continually improving.  For example, in the COSMOS field \\citep{2007ApJS..172....1S}, where we have very deep 30-band photometry from the ultraviolet(GALEX) to $ 5\\mu m$, several photo-z schemes now achieve a precision of $\\sigma_{z} \\sim 0.01(1+z)$, with an outlier fraction (in non-masked areas and defined as a redshift difference greater than $0.15(1+z_{spec})$ of $ <1 \\%$ \\citep{2009ApJ...690.1236I}), at $ I_{AB} < 22.5$ and $ 0.05 < z < 1.4$, where the photo-z can be checked with about over 10,000 spectroscopic redshifts from zCOSMOS. It should be noted in passing that these 30 bands represent a very inhomogeneous data set in terms of point spread function, etc., and so this impressive photo-z performance also demonstrates the feasibility of combining disparate data into homogeneous photometric catalogues. Of course, this outstanding performance in the COSMOS field is unlikely to be achieved over the whole sky for the foreseeable future because of the expense of the required multi-band photometry. Nevertheless, the demonstration of this performance in COSMOS suggests that we have not yet reached any fundamental limit to photo-z performance.  There are a number of different approaches to photo-z estimation that can be broadly distinguished between template-matching and more purely empirical approaches, such as Artificial Neural Networks \\citep{ANN_Lahav}. These have complementary strengths.  Template fitting is based on the observed limited dimensionality of galaxy spectral energy distributions plus an astrophysical knowledge of the effects that can modify them, e.g. the  the redshift itself, and the effects of extinction in our own Galaxy or in the distant galaxy.  The more empirical approaches in essence avoid any such assumptions, which is both a strength and a limitation. Although both approaches have passionate adherents, our own view is that both approaches can normally perform equally well and that both are normally limited by the quality of the available data.  In practice both use elements of the other, e.g. in template fitting, the actual data can be used to adjust the templates and the photometric zero-points, somewhat blurring the distinction.   Finally, it should be noted that both can produce a likelihood distribution in redshift space through the application of priors \\citep{Hyperz, ANN_Lahav, BPZ, EASY}.  In this paper we will base our analysis on a template-fitting algorithm (ZEBRA,~\\citep{2008arXiv0801.3275F}), since we believe its strengths are well suited to the problem in hand. We also note that the impressive real-life performance in COSMOS described above was achieved with two independent template fitting codes (Le Phare, \\citep{2009ApJ...690.1236I} and ZEBRA \\citep{Feldmann2006}).  This improving photo-z performance described above suggests that it may be possible to use the photo-z themselves to construct the N(z), and thus $ \\langle z\\rangle$ for each bin, providing that the systematic uncertainties can be kept below the required level of 0.002(1+z). Some spectroscopic calibration would of course still be required, but the focus of this would be on constructing and characterizing the photo-z algorithm, rather than on constructing the N(z) directly.  This approach would have a number of potential advantages over that discussed by \\cite{Abdalla2008} and others, and summarized above. At the very least, the number of spectra needed may be substantially reduced, although it is unlikely that one would wish to rely entirely on the photo-z and eliminate the spectroscopic conformation completely.  However, to characterize the uncertainties in the photo-z, defined by $ \\sigma_{z}$, to the required level ($ \\sigma_{\\langle z\\rangle}$) we would need of order \\begin{equation} N \\sim  (\\sigma_{z}/\\sigma_{\\langle z \\rangle})^{2} \\end{equation} spectroscopic redshifts, which may be orders of magnitude or more smaller than that implied by equation 1 since $ \\sigma_{z} \\sim 0.05 \\Delta z$.  More importantly, the requirements on completeness and sampling are substantially relaxed since the photo-z characterization is done on individual objects.  As one example, it is relatively easy to simulate the degradation in photo-z performance with noisier photometric data, so the calibrating spectroscopic objects need not necessarily extend all the way down to the photometric limit.  The cosmic variance problem in spectroscopic calibration is eliminated, and the uncertainties arising from Galactic reddening can also be substantially reduced.  \\subsection{Subject of this paper}  The aim of this paper is to explore the performance of a template fitting photo-z code as applied to simulated photometric data of the approximate quality that we may realistically expect for a $ 2\\pi$ sr \\textit{all-sky} ground and space survey within the next decade.  Our emphasis is on both the per object performance and on the potential for recognising and correcting systematic biases, which must be done to a high level if the increased reliance on photo-z is to be possible.  As described in more detail in Section 3, we will assume for definiteness a photometric data set that includes the three-band near infrared photometry that is planned for Euclid itself, plus 5-band grizy photometry similar to that which should be produced by the PanStarrs-1, -2 and -4 projects (hereafter PS-1, -2 and -4)\\footnote{http://pan-starrs.ifa.hawaii.edu}. Examining these three generations of ground-based survey probes a range of depths that can be compared against other future surveys such as DES\\footnote{http://www.darkenergysurvey.org} and LSST\\footnote{http://www.lsst.org}.  We then explore the following four topics that potentially may limit the photo-z performance and their usefulness to construct N(z) and $\\langle z\\rangle$:  \\begin{enumerate} \\item The photo-z performance on individual objects in terms of the r.m.s. scatter (and bias) between the true redshift and the maximum likelihood photo-z, with particular emphasis on how to a priori identify and reject the outliers (\\textit{catastrophic failures}) from their individual photo-z L(z) likelihood distributions. \\item The construction of N(z) for a given set of photo-z selected galaxies, using their photo-z alone, with an emphasis on how to modify their individual likelihood functions L(z) to yield the least biased estimate of N(z) and $\\langle z\\rangle$ for the ensemble. \\item The systematic biases that can enter into the photo-z from an incorrect assumption about the level of foreground Galactic reddening, and how well the photometric data themselves can be used to determine the foreground reddening, both for a set of galaxies at known redshifts, and for those without known redshifts. \\item The effects of the photometric superposition of two galaxies at different redshifts, leading to a mixed spectral energy distribution that may perturb the photo-z, with an emphasis on seeing what happens to the redshift likelihood distribution. An interesting question is whether such composite objects can be recognised photometrically as well as ``morphologically'' from the images. \\end{enumerate}  Our approach is to try to isolate these problems and to explore each in turn with the aim of providing an \\textit{existence proof} that provides a plausible route to achieve the very high photo-z performance that is required for Euclid. In particular, we decided to construct the input photometric catalogues using exactly the same set of approximately ten thousand templates as we subsequently used in the ZEBRA photo-z code.  This may strike some readers as being somewhat circular. However, this approach allows us to eliminate the choice of templates as a variable, or uncertainty, in our analysis.  This is motivated by the exceptional performance (discussed above) that has already been achieved with the same templates coupled with the exquisite observation data in COSMOS.  This strongly suggests to us that the choice of templates is unlikely to be the limiting factor with the degraded photometry that we can realistically expect to have over the whole sky within the timescale of a decade or so.  Although focused on the Euclid cosmology mission, the ideas and results from this paper may be of interest in many other applications that involve photo-z. ",
        "conclusions": "In this work we have investigated a number of issues that could potentially limit the photo-z performance of deep all-sky surveys, and thereby impede the ambitious precision cosmology goals of survey programs such as the proposed ESA Euclid mission.  In each case, we find that, while standard techniques do not get to the required accuracy, simple new approaches can be developed that appear to get to the required performance.  Knowledge of the redshifts enters into weak-lensing analyses at two distinct steps: first, the selection of objects for shape cross-correlation, and second the accurate knowledge of the mean redshifts of these objects. For practical reasons, the first step will likely require the use of photo-z for the foreseeable future.  A major motivation for the current work has been to develop techniques that rely on photo-z also for the second step, bypassing the need for very large and highly statistically complete spectroscopic surveys (e.g. \\citep{Abdalla2008}). This approach thereby avoids the substantial practical difficulties that will be encountered in such spectroscopic surveys from incompleteness and the effects of large scale structure. A separate Appendix explores the latter effects in some detail.  The work is based entirely on simulated photometric catalogues that have been constructed to match the expected performance of three generic ground-based surveys, combined with the expected near-IR imaging photometry from Euclid.  To construct these catalogues, we use the same set of 10,000 templates as used for the template-fitting photo-z program (ZEBRA). This possibly circular approach allows us to remove the choice of templates as a variable.  We believe it is justified at the current time by the very impressive performance of template-fitting photo-z codes applied to the deep multi-band COSMOS photometry, which strongly suggest that the choice of templates will not be a limiting factor at the required level, although further refinement will be desirable.  The analysis is conveniently summarized in terms of the two main requirements on photo-z for precision weak lensing, namely to obtain an r.m.s. precision per object of $\\leq 0.05(1+z)$ and a systematic bias on the mean redshift of a given set of galaxies of $\\leq 0.002(1+z)$. Our main conclusions may be summarized as follows:  \\begin{center} \\begin{itemize} \\item To achieve an r.m.s. photo-z accuracy of $ \\sigma_{z}(z)/(1+z) \\sim 0.05 $ down to $I_{AB} \\leq 24.5$, we need the combination of ground-based photometry with the characteristics of Survey-B (similar to PanStarrs-2 or DES) and the deep all-sky near-infrared survey from Euclid itself.  This performance also requires the implementation of an ``a priori'' rejection scheme (i.e. based on the photometry alone, without knowledge of the actual redshifts of any galaxies) that rejects 13$ \\%$ of the galaxies and reduces the fraction of $ 5\\sigma$ outliers to below $ f_{cat} < 0.25\\%$. There is a trade off between the rejection of outliers and the loss of ``innocent'' galaxies with usable photo-z. Deeper photometry improves both the statistical accuracy of the photo-z and reduces the wastage in eliminating the catastrophic failures, and the combination of survey-C with Euclid near-infrared photometry achieves $ \\sigma_{z}(z)/(1+z) \\sim 0.04 $ after 9$ \\%$ rejection.  \\item A good way to determine the actual $N(z)$ of a set of galaxies in a given photo-z selected redshift bin is to sum their individual likelihood functions.  We find that the $\\sum L(z)$ function represents well both the wings of the $N(z)$ and the remaining catastrophic failures. outliers.  However, to reach the required performance on the mean of the redshifts $| \\Delta_{\\langle z\\rangle} | \\leq 0.002(1+z)$ with the Survey-B, or deeper survey-C, combination (together with Euclid infrared photometry), we had to implement a modification scheme on the individual $L(z)$. This is based on the spectroscopic measurement of redshifts for a rather small number of galaxies (less than 1000) with relaxed requirements on statistical completeness (and no dependence on Large Scale Structure in the spectroscopic survey fields). We then require that the distribution of the actual redshifts within the probability space that is defined by the individual $L(z)$ should be flat across the sample as a whole.  This should be true of any set of galaxies, leading to relaxed requirements on sampling of the redshift survey.  This approach is similar to the application of a Bayesian prior on the redshifts, but is performed in probability space.  Although it cannot be rigourously justified, it is found to work well in practice.  \\item  We find that uncertainties in foreground Galactic reddening can have a serious effect in perturbing the photo-z, with a net sign that varies erratically with redshift. However, we also find that such errors in $A_{v}$ can be identified internally from the photometric data of galaxies, either with or without spectroscopic redshifts. This procedure works best for galaxies with relatively high S/N photometry $ I_{AB} \\leq 22$. The required number of galaxies suggests that a reddening map on the scales of 0.1 deg$^{2}$ can be internally constructed from the data on galaxies without known redshifts, or from a few hundred galaxies with spectroscopic redshifts.  \\item We also explored the effect of blended objects. The photo-z of the composite spectral energy distribution is a good representation of the redshift of the brighter object as long as the magnitude difference is large, i.e. $ \\Delta I_{AB} > 2.$ When the galaxies are more similar in brightness, $ \\Delta I_{AB} < 2$ there is a wide-range of behaviour. In some cases, multi-modal likelihood functions appear, while in others there is a sharp transition from one redshift to another, sometimes with a local maximum at a third, completely spurious redshift.  In still others, the likelihood function smoothly transitions between the two redshifts with a single peak at an intermediate redshift.  Our conclusion is that composite objects with $ \\Delta I_{AB} < 2$ should be recognized morphologically from imaging data and removed from the photo-z analysis.  \\end{itemize} \\end{center}  The general conclusion of our study is that while reaching the photo-z performance required for weak-lensing surveys such as Euclid will not be trivial, the implementation of new techniques, coupled with internal calibration of e.g. foreground reddening from the photometric data itself, will allow the required performance to be attained.  If more reliance is placed on the photo-z themselves, then this leads to a significant simplification of the otherwise challenging requirement for spectroscopic calibration of large scale photometric surveys."
    },
    "0910/0910.5359_arXiv.txt": {
        "abstract": "The description of high-energy hadronic interactions plays an important role in the (astrophysical) interpretation of air shower data.  The parameter space important for the development of air showers (energy and kinematical range) extends well beyond todays accelerator capabilities.  Therefore, accurate measurements of air showers are used to constrain modern models to describe high-energy hadronic interactions.  The results obtained are complementary to information gained at accelerators and add to our understanding of high-energy hadronic interactions. ",
        "introduction": "The understanding and modeling of extensive air showers (particle cascades in the atmosphere) brings together the particle physics and astroparticle physics communities. To strengthen the connections and the scientific exchange between those communities is very fruitful for both sides and yields complementary information on the understanding of high-energy hadronic interactions.  When high-energy cosmic rays impinge onto the atmosphere they initiate cascades of secondary particles -- the extensive air showers.  Observations of air showers are used to indirectly infer the properties of cosmic rays at energies exceeding $10^{14}$~eV.  The interpretation of air shower data faces a twofold challenge: the (exact) mass composition of cosmic rays is not known at those energies and, additionally, the properties of high-energy interactions taking place in air showers are partly unknown.  Direct measurements of cosmic rays (fully ionized atomic nuclei) at energies below $10^{14}$~eV indicate that they are mostly composed of elements from hydrogen (protons) up to iron \\cite{cospar06,behreview}. The abundance of heavier elements is significantly smaller. Hence, in the following, we assume that cosmic rays comprise elements from hydrogen to iron.  We will focus on results from the KASCADE experiment \\cite{kascadenim}, one of the most advanced air shower detectors in the energy range around $10^{15}$~eV. It has a unique set-up which allows to measure simultaneously the electromagnetic, muonic, and hadronic shower components.  This is in particular valuable to test the consistency of hadronic interaction models.  Since about a decade \\cite{Horandel:1998br,wwtestjpg} systematic checks of interaction models are performed with air shower data and the most stringent constraints on interaction models, derived from air shower data have been obtained with KASCADE measurements.  KASCADE consists of several detector systems. A $200 \\times 200$~m$^2$ array of 252 detector stations, equipped with scintillation counters, measures the electromagnetic and, below a lead/iron shielding, the muonic parts of air showers.  An iron sampling calorimeter of $16 \\times 20$~m$^2$ area detects hadronic particles \\cite{kalonim}. It has been calibrated with a test beam at the SPS at CERN up to 350~GeV particle energy \\cite{kalocern}.  For a detailed description of the reconstruction algorithms see \\rref{kascadelateral}. ",
        "conclusions": ""
    },
    "0910/0910.4396_arXiv.txt": {
        "abstract": "We study spectral distortions of diffuse ultra-high energy (UHE) neutrino flavour fluxes resulting due to physics beyond the Standard Model (SM). Even large spectral differences between flavours at the source are massaged into a common shape at earth by SM oscillations, thus, any significant observed spectral differences are an indicator of new physics present in the oscillation probability during propagation. Lorentz symmetry violation (LV) and Neutrino decay are examples, and result in significant distortion of the fluxes and of the well-known bounds on them, which may allow UHE detectors to probe LV parameters, lifetimes and the mass hierarchy over a broad range. ",
        "introduction": "Introduction} The neutrino sky spans about twenty five orders of magnitude in energy, potentially offering the possibility of probing the universe at widely disparate energy scales. The high end ($10^{11}-10^{12} $GeV) of this remarkably broad band in energy is set by: a) GZK neutrinos \\cite{PhysRevLett.16.748,Zatsepin:1966jv}, which originate in the interactions of the highest energy cosmic rays with the cosmic microwave photon background and b) neutrinos from the most energetic astrophysical objects observed in the universe, \\textit{i.e.} active galactic nuclei (AGNs) and Gamma-Ray bursts (GRBs). Detection, still in the future, presents both considerable opportunities and formidable challenges. In particular, at the very highest energies ($10^5$ GeV and above) which are the focus of this paper, the tiny fluxes that arrive at earth require detectors that combine the capability to monitor very large detection volumes with innovative techniques (for reviews, see e.g. \\cite{Hoffman:2008yu} and \\cite{Halzen:2007sz}). Examples of such detectors are AMANDA \\cite{Baret:2008zz}, ICECUBE \\cite{Halzen:2009tz}, BAIKAL \\cite{Aynutdinov:2009zz}, ANTARES \\cite{Margiotta:2009zz}, RICE \\cite{1742-6596-81-1-012008} and ANITA \\cite{1742-6596-136-2-022052}.  A compelling motivation for exploring UHE  neutrino astronomy is the fact that the origin of cosmic rays (CR) beyond the ``knee\" ($10^6$ GeV) remains a mystery many decades after their discovery. Additionally, CR with energies in excess of $10^{11}$ GeV have been observed\\cite{Roth:2009zz,Abbasi:2007sv}, signalling the presence of astrophysical particle accelerators of unprecedentedly high energies. If protons as well as neutrons are accelerated at these sites in addition to electrons, standard particle physics predicts correlated fluxes of neutrons and neutrinos which escape from the confining magnetic field of the source, while protons and electrons stay trapped. Generically, the electrons lose energy rapidly via synchrotron radiation. These radiated photons provide a target for the accelerated protons, which results in the production of pions, muons and ultimately, neutrinos in the ratio ${\\nu_e:\\nu_\\mu:\\nu_\\tau = 1:2:0}$. The detection and study of ultra-high energy (UHE) neutrinos is thus a probe of the origin of CR and the physics of UHE astrophysical accelerators.  As is well understood due to a wealth of data from solar, atmospheric, reactor and accelerator  experiments, neutrinos have mass and oscillate into one another. These  experiments have led to determinations, to various degrees of accuracy, of the neutrino mass (squared) differences $ \\Delta m^2_{ij} $ between mass eigenstates $ i $ and $ j $, and the mixing angles $ U_{\\alpha i} $ \\cite{Schwetz:2009zz}. The latter characterize the overlap of neutrino mass (or propagation) eigenstates (denoted by $\\nu_i, i=1,2,3$) and the interaction eigenstates (denoted by $\\nu _\\alpha, \\alpha=e,\\mu,\\tau$). Over the very large distances traversed by neutrinos from the most energetic extragalactic hadronic accelerators, the initial source ratios ${\\nu_e^s:\\nu_\\mu^s:\\nu_\\tau^s = 1:2:0}$ will, due to (vacuum) oscillations, transmute, by the time they reach a terrestrial detector, into $\\nu_e^d:\\nu_\\mu^d:\\nu_\\tau^d = 1:1:1$ \\cite{Learned:1994wg,Athar:2000yw}.  A series of papers \\cite{Beacom:2002vi,Beacom:2003nh,Barenboim:2003jm,Keranen:2003xd,Beacom:2003eu,Hooper:2004xr,Hooper:2005jp, Meloni:2006gv,Xing:2008fg,Esmaili:2009dz} over the past few years have demonstrated that if the flavour ratios $\\nu_e^d:\\nu_\\mu^d:\\nu_\\tau^d$ detected by extant and upcoming neutrino telescopes were to deviate significantly from this democratic prediction, then important conclusions about physics beyond the Standard Model and neutrino oscillation parameters may consequently be inferred. In addition, deviations of these  measured ratios have been shown to be sensitive to  neutrino oscillation parameters \\cite{Farzan:2002ct, Beacom:2003zg, Serpico:2005sz, Xing:2006xd, Rodejohann:2006qq, Winter:2006ce, Awasthi:2007az, Lipari:2007su, Blum:2007ie, Pakvasa:2007dc, Choubey:2008di, Esmaili:2009dz} (\\textit{e.g.}~the  mixing angles  and the Dirac CP violating phase).  In this paper we  study the spectral distortions in the {\\it diffuse} ({\\it i.e.} integrated over source distribution and redshift) UHE neutrino flux as a probe for the  effects of new physics. For specificity, we focus on AGN fluxes, and use, as a convenient bench-mark, the well-known upper bounds first derived by Waxman and Bahcall (WB) \\cite{Waxman:1998yy} and later by Mannheim, Protheroe and Rachen (MPR) \\cite{Mannheim:1998wp} on such fluxes for both neutron-transparent  and neutron-opaque sources (or, equivalently, sources that are \\emph{optically thin} and \\emph{optically thick}, respectively, to the emission of neutrons). In particular we focus on the upper bounds to the diffuse neutrino flux from hadronic photoproduction in AGN's derived in \\cite{Mannheim:1998wp} using the experimental upper limit on cosmic ray protons. All distortions in the fluxes are, as would be expected, transmitted to the upper bounds, thus providing a convenient way of representing and studying them.  Prior to this, we first demonstrate (in Sec.~\\ref{sec:spectral-averaging}) that the usual (SM) neutrino oscillations not only tend to equilibrate widely differing source flux magnitudes between flavours, but also massage them into a common spectral shape, as one would intuitively expect. Thus observed relative spectral distortions among flavours are a probe  of new physics present in the propagation equation. To demonstrate our approach we then  focus on two specific examples, \\textit{i.e.,} \\begin{inparaenum}[\\itshape a\\upshape)] \\item Lorentz violation (LV), and \\item Incomplete decay of neutrinos \\end{inparaenum} in the neutrino sector. Our method can straightforwardly be applied to other new physics scenarios and our results translated into bounds on muon track versus shower event rates\\footnote{ These count the sum of \\begin{inparaenum}[\\itshape a\\upshape)] \\item neutral current (NC) events of all flavours, \\item electron neutrino charged current (CC) events at all energies, and \\item $\\nu_\\tau$ induced CC events at energies below $\\leq$ 1 PeV ($10^6$ GeV), \\end{inparaenum} whereas muon track events arise from $\\nu_\\mu$ induced muons born in CC interactions.} for UHE experiments.  In our calculations we use the current best-fit values of neutrino mixing paramenters as given in \\cite{Maltoni:2008ka}. The mass squared differences and mixing angles are \\begin{gather*} \\begin{split} \\Delta m_{21}^2 &= 7.65 \\times 10^{-5} \\text{ eV}^2 \\\\ \\Delta m_{31}^2 &= \\pm 2.40 \\times 10^{-3} \\text{ eV}^2 \\end{split} \\\\ \\sin^2(\\theta_{12})=0.321,\\;\\; \\sin^2(\\theta_{23})=0.47,\\;\\; \\sin^2(\\theta_{13})=0.003. \\end{gather*} ",
        "conclusions": "The detection of UHE neutrinos is imminent. Several detectors will progressively sharpen their capabilities to detect neutrino flavours, beginning with ICECUBE's ability to separate muon tracks from shower events. We have shown that spectral changes in diffuse UHE neutrino fluxes of different flavours are probes of new physics entering the oscillation probability. For specificity we have used AGN sources, and calculated the changes induced in the well-known MPR bounds on both neutron-opaque and neutron-transparent sources. Our calculations can, in a straightforward manner, be repeated  for other sources, or represented in terms of the WB bounds or in terms of actual fluxes and  event rates.  We have shown that diffuse UHE fluxes are sensitive to the presence, over a relatively broad range, of tiny LV terms in the effective oscillation hamiltonian, which signal important breakdowns of pillars of presently known physics. Additionally, in the specific example of neutrino decay,  UHE detectors can probe the unexplored range $ 10^{-3} \\text{ s/eV }$  $< \\tau/m < 10^4 \\text{ s/eV } $ via spectral differences in flavours, and thus  extend our knowledge of this important parameter. Flavour spectra in decay scenarios also differ depending on the mass hierarchy, and  provide a potent tool to probe it.  However, it must be kept in mind that the ability to distinguish between flavours in most present-day UHE detectors is as yet not firmly established. At present, detectors are capable of separating between shower events and muon-tracks which include a sum of contributions from each of the flavours rather than dinstinguishing between the individual contributions themselves. Further, the energy resolutions of present or near-future large-volume detectors like the IceCube are not very high. Consequently, deviation of the flavour spectra from the standard predictions has to be significantly large, both in magnitude and shape, for detectability in any of these detectors. Our calculations show that it is indeed possible for such large deviation to occur in the diffuse fluxes of the three flavours for suitable values of the parameters in the case of Lorentz violation as well as decay with inverted hierarchy.  Other physics scenarios to which this method may be effectively applied are the existence of pseudo-dirac states, CP violation and quantum decoherence. Although the detection of the effects discussed here will undoubtedly be challenging, given the fundamental nature of the physics to be probed, it would certainly be worthwhile.  \\textbf{\\textit{Acknowledgment:}} RG would like to thank Nicole Bell, Maury Goodman, Francis Halzen, Chris Quigg, Georg Raffelt, Subir Sarkar and Alexei Smirnov  for helpful discussions, and the theory groups at Fermilab and Brookhaven National Lab for hospitality while this work was in progress."
    },
    "0910/0910.4958_arXiv.txt": {
        "abstract": "{We present new measurements of the large-scale bulk flows of galaxy clusters based on 5-year WMAP data and a significantly expanded X-ray cluster catalogue. Our method probes the flow via measurements of the kinematic Sunyaev-Zeldovich (SZ) effect produced by the hot gas in moving clusters. It computes the dipole in the cosmic microwave background (CMB) data at cluster pixels, which preserves the SZ component while integrating down other contributions. Our improved catalog of over 1,000 clusters enables us to further investigate possible systematic effects and, thanks to a higher median cluster redshift, allows us to measure the bulk flow to larger scales.  We present a corrected error treatment and demonstrate that the more X-ray luminous clusters, while fewer in number, have much larger optical depth, resulting in a higher dipole and thus a more accurate flow measurement. This results in the observed correlation of the dipole derived at the aperture of zero monopole with the monopole measured over the cluster central regions. This correlation is expected if the dipole is produced by the SZ effect and cannot be caused by unidentified systematics (or primary cosmic microwave background anisotropies). We measure that the flow is consistent with approximately constant velocity out to at least $\\simeq$800 Mpc. The significance of the measured signal peaks around 500 $h_{\\rm 70}^{-1}$Mpc, most likely because the contribution from more distant clusters becomes progressively more diluted by the WMAP beam. We can, however, at present not rule out either that these more distant clusters simply contribute less to the overall motion.} ",
        "introduction": "KABKE1,2 and this work utilize a method of Kashlinsky \\& Atrio-Barandela (2000, hereafter KA-B), which measures CMB dipole at the locations of X-ray clusters. When averaged over many isotropically distributed clusters moving at a significant bulk flow with respect to the CMB, the kinematic term dominates the SZ signal, thereby enabling a measurement of $V_{\\rm bulk}$ over that distance. In Atrio-Barandela et al (2008, hereafter AKKE) we demonstrated that 1) the thermal SZ (TSZ) signal from clusters extends well beyond the measured X-ray extent $\\Theta_{\\rm X}$, and 2) the intra-cluster gas distribution is well approximated by the NFW profile (Navarro et al 1996) expected for dark matter in a $\\Lambda$CDM model. The temperature $T_X$ of hot gas distributed according to these profiles decreases significantly from the cluster cores to the cluster outskirts (Komatsu \\& Seljak 2001), consistent with current measurements (Pratt et al 2007) and numerical simulations (e.g. Borgani et al 2004). Consequently, the monopole produced by the TSZ component ($\\propto \\tau T_X$) decreases, as we increase the cluster aperture, whereas the dipole due to the kinematic SZ (KSZ) component ($\\propto \\tau$, the optical depth due to Thomson scattering) remains measurable out to the aperture where we still detect the TSZ decrement in unfiltered maps (KABKE2). As in KABKE1,2 our dipole coefficients are normalized such that the dipole power $C_{1}$ due to a coherent motion at velocity $V_{\\rm bulk}$ is $C_{1,{\\rm kin}}= T_{\\rm CMB}^2 \\langle \\tau \\rangle^2 V_{\\rm bulk}^2/c^2$, where $T_{\\rm CMB} =2.725$K. We adopt $\\Omega_{\\rm total}=1, \\Omega_\\Lambda=0.7, H_0=70h_{70}$ km/sec/Mpc.  To improve upon the all-sky cluster catalogue of Kocevski \\& Ebeling (2006) used by KABKE1,2, we have screened the ROSAT Bright-Source Catalogue (Voges et al.\\ 1999) using the same X-ray selection criteria (including a nominal flux limit of $1\\times 10^ {-12}$ erg/sec, 0.1--2.4 keV) as employed during the Massive Cluster Survey (MACS, Ebeling et al.\\ 2001), as well as the same optical follow-up strategy. Unlike MACS, we apply neither a declination nor a redshift limit though, thereby creating an all-sky list of cluster candidates that extends to three times fainter X-ray fluxes than in KABKE1,2. Optical follow-up observations of clusters identified in this manner and lacking spectroscopic redshifts in the literature (and MACS) are well underway using telescopes on Mauna-Kea/Hawai`i and La-Silla/Chile. As a result, our interim all-sky cluster catalogue currently comprises in excess of 1,400 X-ray selected clusters, {\\it all} of them with spectroscopic redshifts. X-ray properties of all clusters (most importantly total luminosities and central electron densities) are computed as before (KABKE2). Within the same $z$-range ($z\\la0.25$) as our previous study, our new catalog comprises 1,174 clusters outside the KP0 CMB mask. To eliminate low-mass galaxy groups we require that clusters feature $L_X\\geq 2 \\times 10^{43}$ erg/sec; 985 systems meet this criterion. This sample represents a significant improvement over the one in KABKE1,2 largely because of the substantial increase ($z_{\\rm median}<0.1$ in KABKE1,2 vs $z_{\\rm median}\\simeq 0.2$ below) in median cluster redshift and the higher fraction of intrinsically very X-ray luminous systems (Fig. 1a).  We applied this new cluster sample to the 5-year WMAP CMB data, processed as described in detail in KABKE1,2. The all-sky dipole in the foreground-cleaned maps from http://www.lambda.gsfc.nasa.gov was removed, and standard CMB masking was applied. We also need to remove the primary CMB fluctuations, produced at last scattering, as they are highly correlated and would contribute significantly to the measured dipole. To this end, all maps were filtered as in KABKE1,2 with a filter that minimizes $\\langle (\\delta T - \\delta T_{\\Lambda{\\rm CDM}})^2\\rangle$. The error budget associated with our filtering is discussed by Atrio-Barandela et al (2010, hereafter AKEKE).  Measurement errors were computed as in KABKE1,2 with one important correction. Although the filtering removes much of the CMB fluctuations, a residual component remains, due to cosmic variance and imperfections of the theoretical model. Since this residual is common to all WMAP bands, a component of the errors is correlated between the various DA maps. We address this issue in the following manner. For each of the eight differential assemblies (DAs) we simulate 4,000 realizations with $N_{\\rm cl}=$100--1,000 randomly selected pseudo-clusters outside of the cluster pixels and the CMB mask. In each realization we select the {\\it same} pseudo-clusters for all DAs and evaluate the {\\it mean} monopole and dipole averaged over all DAs: $\\bar{a}_0, \\bar{a}_{1m}$. For the 4,000 realizations at each $N_{\\rm cl}$ we compute the mean and dispersion of $\\bar{a}_0, \\bar{a}_{1m}$ over all the realizations. The distributions have zero mean and their dispersion gives errors which scale as $N_{\\rm cl}^{-1/2}$. We find to good accuracy that the distribution of the simulated dipoles (and monopoles) is Gaussian and the errors on each of the averaged dipole components are $\\sigma_{1m} \\simeq 15 \\sqrt{3/N_{\\rm cl}}\\mu$K and on the monopole $\\sigma_0 \\simeq 15 \\sqrt{1/N_{\\rm cl}} \\mu$K as explained in great detail in AKEKE; the errors of the $x/z$ dipole component are slightly larger/smaller because of the Galactic mask (Fig. 1b). ",
        "conclusions": "We find a high likelihood of the existence of a coherent bulk flow extending to at least $z\\simeq 0.2$ with an amplitude and in a direction which are in good agreement with our earlier measurements. Our result constitutes a significant improvement in that it extends our previous work to approximately twice the distance accessible to KABKE1,2, supporting their hypothesis that the flow likely extends across much (or all) of the Hubble volume. The flow's axis is also consistent with earlier measurements of the local cluster dipole (Kocevski et al.\\ 2004) as well as with independent measurements of bulk flows on smaller scales by Watkins et al (2009). The velocity reported there is smaller than the numbers in Table 1, although the two amplitudes agree at $< 2$-$\\sigma$ level. Agreement between the two sets of central values would require $\\sqrt{\\langle C_{1,100}\\rangle} \\sim$0.4-0.5$\\mu$K, or a reduction by a factor of $\\sim2$ from unfiltered values for NFW cluster profiles. Feldman et al (2009) extend the Watkins et al analysis and find that the absence of shear in their flow at $\\la$50-100Mpc is consistent with the KABKE1 suggestion of the attractor at superhorizon distances.  Fig. \\ref{fig:f2} displays the results obtained in this study compared to expectations from the concordance $\\Lambda$CDM model for 95\\% of cosmic observers. These results cast doubt on the notion that gravitational instability from the observed mass distribution is the sole -- or even dominant  -- cause of the detected motion. If the current picture is confirmed, it will have profound implications for our understanding of the global structure of space-time and our Universe's place in it.  We acknowledge NASA NNG04G089G/09-ADP09-0050 and FIS2006-05319/GR-234 grants from Spanish Ministerio de Educaci\\'on y Ciencia/Junta de Castilla y Le\\'on.   {\\bf REFERENCES}\\\\ Afshordi, N. Geshnizjahi, G. \\& Khoury, J. 2009, JCAP, 8, 30\\\\ Atrio-Barandela, F., Kashlinsky, A., Kocevski, D. \\& Ebeling, H. 2008, Ap.J. (Letters), 675, L57 (AKKE)\\\\ Atrio-Barandela, F., Kashlinsky, A., Ebeling, H., Kocevski, D. \\& Edge, A. 2010, ApJ, submitted, arXiv:1001.1261 (AKEKE)\\\\ Borgani, S. et al  2004, Mon. Not. R. Astron. Soc., 348, 1078\\\\ Carrol, S. et al 2008, arxiv:0811.1086\\\\ Ebeling, H.; Edge, A. C.; Henry, J. P., 2001, ApJ, 553, 668\\\\ Feldman, H., Watkins, R. \\& Hudson, M.J. 2009, arXiv:0911.5516v1\\\\ Gorski, K. et al 2005, Astrophys. J., 622, 759\\\\ Grishchuk, L. P. 1992, Phys. Rev. D 45, 4717\\\\ Kashlinsky, A., Tkachev, I., Frieman, J.  1994, Phys. Rev. Lett., 73, 1582\\\\ Kashlinsky, A. \\& Atrio-Barandela, F.  2000, Astrophys. J., 536, L67 (KA-B)\\\\ Kashlinsky, A., Atrio-Barandela, F., Kocevski, D. \\& Ebeling, H. 2008, Ap.J., 686, L49 (KABKE1)\\\\ Kashlinsky, A., Atrio-Barandela, F., Kocevski, D. \\& Ebeling, H. 2009, Ap.J., 691, 1479 (KABKE2)\\\\ Kashlinsky, A. \\& Jones, B.J.T. 1991, Nature, 349, 753\\\\ Khoury, J. \\& Wyman, M. 2009, PhRevD, 80, 064023\\\\ Kocevski, D.D., Mullis, C.R., \\& Ebeling, H. 2004, Astrophys. J., 608, 721\\\\ Kocevski, D.D. \\& Ebeling, H.  2006, Astrophys. J., 645, 1043\\\\ Komatsu, E. \\& Seljak, U. 2001, Mon. Not. R. Astron. Soc., 327, 1353\\\\ Mersini-Houghton, L. \\& Holman, R. 2009, JCAP,  2, 6\\\\ Navarro, J.F., Frenk, C.S. \\& White, S.D.M. 1996, Astrophys. J., 462, 563\\\\ Pratt, G.  et al 2007, Astron. Astrophys. 461, 71\\\\ Turner, M. S. 1991, Phys.Rev., 44, 3737\\\\ Watkins, R., Feldman, H. A. \\& Hudson, M. J. 2009, MNRAS, 392, 743\\\\ Voges, W. et al 1999, A\\&A, 349, 389\\\\  \\clearpage \\thispagestyle{empty} \\begin{deluxetable}{l c c c c c | c | c c | c c c c c} \\tablewidth{0pt} \\tabletypesize{\\scriptsize} \\rotate \\tablecaption{RESULTS  \\label{table} } \\tablehead{ \\colhead{(1)} & \\colhead{(2)} & \\colhead{(3)} & \\colhead{(4)} & \\colhead{(5)} & \\colhead{(6)} & \\colhead{(7)} & \\colhead{(8)} & \\colhead{ } & \\colhead{ } & \\colhead{ } & \\colhead{(9)}  & \\colhead{ }  & \\colhead{ }  \\\\ } \\startdata $z\\leq$ & $L_X$-bin & $N_{\\rm cl}$ & $z_{\\rm mean}/z_{\\rm median}$  & $\\bar{a}_{1,x}, \\bar{a}_{1,y}, \\bar{a}_{1,z}$ & $\\sqrt{C_1}$ & $\\langle\\tau_0\\rangle$ & $\\sqrt{C_{1,100}}$ & ($\\mu$K) & & & $\\bar{a}_0$ & ($\\mu$K) & \\\\ & $10^{44}$ erg/s &  &  & $\\mu$K & $\\mu$K & $\\times10^{-3}$ & 5$^\\prime$ & $\\Theta_{\\rm X}$ & $10^\\prime$ & $15^\\prime$ & $20^\\prime$ & $25^\\prime$ & $30^\\prime$\\\\ \\hline 0.12$^{*}$ & 0.2--0.5 & 142 & 0.061/0.060 & $-4.2\\pm2.7, -0.7\\pm2.3, 0.5\\pm2.3$  & $4.3\\pm2.7$ & 2.8 & 0.2301 &  0.1942 &  --2.8 & 0.1  & -- & -- & -- \\\\ 0.12 & 0.5--1 & 194 & 0.081/0.082 & $-2.7\\pm2.3, -2.3\\pm2.0, 1.4\\pm2.0$  & $3.9\\pm2.2$ & 3.5 & 0.2989 & 0.2561 & --2.4 & --1.2  & --0.1 & 0.6 & 0.8 \\\\ 0.12 & $>1$ & 180 & 0.083/0.086 & $4.9 \\pm 2.4, -4.5 \\pm 2.1, 1.5\\pm 2.0$ & $6.8\\pm2.2$ & 5.4 & 0.4610 & 0.3496  & --11.1 & --6.5  & --3.1 & --0.8 & 0.5 \\\\ \\hline \\\\ \\multicolumn{14}{l}{\\vspace{0.1in} $d\\sim 250-370h_{70}^{-1}$Mpc; $(V_x,V_y,V_z)= (174 \\pm 407, -849 \\pm 351 , 348 \\pm 342 )\\times\\frac{+0.3\\mu K}{\\sqrt{\\langle{C}_{1,100}\\rangle}}$ km/sec; $V_{\\rm Bulk}=(934 \\pm 352)\\times\\frac{0.3\\mu K}{\\sqrt{\\langle{C}_{1,100}\\rangle}}$ km/sec; $(l_0,b_0)= (282\\pm 34, 22 \\pm 20)^\\circ$} \\\\ \\hline 0.16 & 0.5--1 & 226 & 0.089/0.087 & $-1.5\\pm2.2, -0.6 \\pm1.9, 2.1 \\pm 1.8$  & $2.7\\pm1.9$ & 3.5 & 0.2843 & 0.2363 & --2.8 & --1.8  & --0.6 & 0.2 & 1.6 \\\\ 0.16 & 1--2 & 191 & 0.106/0.107 & $1.9\\pm2.3, -2.8 \\pm 2.0, -0.5 \\pm 2.0$  & $4.1\\pm2.2$ & 4.4 & 0.3480 & 0.2894 & --4.9 & --1.4  & 0.4 & 1.3 & 1.8 \\\\ 0.16 & $>2$ & 130 & 0.115/0.125 & $4.2\\pm2.8, -8.0\\pm2.4, 4.9\\pm2.4$  & $10.3\\pm2.5$ & 6.8 & 0.4930 & 0.4238 & --11.7 & --7.1  & --2.9 & --0.3 & 0.8 \\\\ \\hline \\\\ \\multicolumn{14}{l}{$^{(a)}$ $d\\sim 370-540 h_{70}^{-1}$Mpc; $(V_x,V_y,V_z) = (410 \\pm 379 , -1,012 \\pm 326, 566 \\pm 319 )\\times\\frac{+0.3\\mu K}{\\sqrt{\\langle{C}_{1,100}\\rangle}}$ km/sec; $V_{\\rm Bulk}=(1,230\\pm 331 )\\times\\frac{0.3\\mu K}{\\sqrt{\\langle{C}_{1,100}\\rangle}}$ km/sec; $(l_0,b_0)=( 292\\pm21, 27 \\pm 15)^\\circ$} \\\\ \\multicolumn{14}{l}{\\vspace{0.1in} $^{(b)}$ $d\\sim 370-540 h_{70}^{-1}$Mpc; $(V_x,V_y,V_z)= (428 \\pm 375 , -1,029 \\pm 323, 575 \\pm 316 )\\times\\frac{+0.3\\mu K}{\\sqrt{\\langle{C}_{1,100}\\rangle}}$ km/sec; $V_{\\rm Bulk}=(1,254\\pm 328 )\\times\\frac{+0.3\\mu K}{\\sqrt{\\langle{C}_{1,100}\\rangle}}$ km/sec; $(l_0,b_0)=(293\\pm 20, 27\\pm 15)^\\circ$} \\\\ \\hline 0.20 & 0.5--1 & 238 & 0.093/0.089 & $-2.5\\pm2.1, -1.3 \\pm 1.8, 1.0 \\pm 1.8$ & $3.0\\pm2.0$ & 3.5 & 0.2828 & 0.2390 & --2.9 & --2.2  & --1.1 & --0.3 & --0.2 \\\\ 0.20 & 1--2 & 248 & 0.122/0.123 & $0.1\\pm2.0, -1.8 \\pm 1.8, -0.3\\pm 1.7$ & $1.8\\pm1.8$ & 4.4 & 0.3231 & 0.2835 & --5.1 & --1.8  & --0.3 & 0.5 & 1.0 \\\\ 0.20 & $>2$ & 208 & 0.140/0.151 & $3.6\\pm2.2, -5.8 \\pm1.9, 4.5\\pm1.9$ & $8.1\\pm2.0$ & 6.6 & 0.4644 & 0.4218 & --9.3 & --5.5  & --1.9 & 0.4 & 1.1 \\\\ \\hline \\\\ \\multicolumn{14}{l}{$^{(a)}$ $d\\sim 380-650 h_{70}^{-1}$Mpc; $(V_x,V_y,V_z)= (213 \\pm 341, -872\\pm 294, 529 \\pm 287)\\times\\frac{+0.3\\mu K}{\\sqrt{\\langle{C}_{1,100}\\rangle}}$ km/sec; $V_{\\rm Bulk}=(1,042 \\pm 295)\\times\\frac{0.3\\mu K}{\\sqrt{\\langle{C}_{1,100}\\rangle}}$ km/sec; $(l_0,b_0)=(284\\pm 24, 30 \\pm 16)^\\circ$} \\\\ \\multicolumn{14}{l}{\\vspace{0.1in} $^{(b)}$ $d\\sim 380-650 h_{70}^{-1}$Mpc; $(V_x,V_y,V_z)= (248 \\pm 337, -880\\pm 291, 538 \\pm 284)\\times\\frac{0.3\\mu K}{\\sqrt{\\langle{C}_{1,100}\\rangle}}$ km/sec; $V_{\\rm Bulk}=(1,061 \\pm 292)\\times\\frac{0.3\\mu K}{\\sqrt{\\langle{C}_{1,100}\\rangle}}$ km/sec; $(l_0,b_0)=(286\\pm 23, 30 \\pm 15)^\\circ$} \\\\ \\hline 0.25 & 0.5--1 & 240 & 0.094/0.090 & $-2.3\\pm2.1, -1.1 \\pm 1.8, 0.9 \\pm 1.8$ & $2.7\\pm2.0$ & 3.5 & 0.2848 & 0.2444 & --2.8 & --2.1  & --1.0 & --0.3 & --0.1 \\\\ 0.25 & 1--2 & 276 & 0.133/0.133 & $-0.2 \\pm 2.0, -1.4\\pm 1.7, 0.7 \\pm 1.6$ & $1.6\\pm1.7$ & 4.4 & 0.3162 & 0.2806 & --5.8 & --2.3  & --0.8 & -0.1 & 0.3 \\\\ 0.25 & $>2$ & 322 & 0.169/0.176 & $3.7\\pm1.8,-4.1\\pm1.5,4.1\\pm1.5$ &  $6.9\\pm1.6$  & 6.6 & 0.4434 & 0.4160 & --6.9 & --4.6  & --2.3 & --0.6 & 0.2 \\\\ \\hline \\\\ \\multicolumn{14}{l}{$^{(a)}$ $d\\sim 385-755 h_{70}^{-1}$Mpc; $(V_x,V_y,V_z)= (313\\pm 308, -707\\pm 265, 643\\pm 259)\\times\\frac{+0.3\\mu K}{\\sqrt{{\\langle{C}}_{1,100}\\rangle}}$ km/sec; $V_{\\rm Bulk}=(1,005\\pm267)\\times\\frac{0.3\\mu K}{\\sqrt{\\langle{C}_{1,100}\\rangle}}$ km/sec; $(l_0,b_0)=(296 \\pm 29, 39 \\pm 15)^\\circ$} \\\\ \\multicolumn{14}{l}{\\vspace{0.1in} $^{(b)}$ $d\\sim 385-755 h_{70}^{-1}$Mpc; $(V_x,V_y,V_z)= (352\\pm 304, -713\\pm 262, 652\\pm 256)\\times\\frac{+0.3\\mu K}{\\sqrt{{\\langle{C}}_{1,100}\\rangle}}$ km/sec; $V_{\\rm Bulk}=(1,028\\pm265)\\times\\frac{0.3\\mu K}{\\sqrt{\\langle{C}_{1,100}\\rangle}}$ km/sec; $(l_0,b_0)=(296 \\pm 28, 39 \\pm 14)^\\circ$} \\\\ \\hline \\enddata \\tablecomments{ } \\end{deluxetable} \\clearpage \\begin{figure} \\plotone{f1.eps} \\caption{ \\small{ }} \\label{fig:f1} \\end{figure}  \\clearpage  \\begin{figure} \\plotone{f2.eps} \\caption{ \\small{ }} \\label{fig:f2} \\end{figure}   \\clearpage"
    },
    "0910/0910.1099_arXiv.txt": {
        "abstract": "We present observations at 1.2~mm with MAMBO-II of a sample of $z\\grtsim2$ radio-intermediate obscured quasars, as well as CO observations of two sources with the Plateau de Bure Interferometer. The typical rms noise achieved by the MAMBO observations is 0.55~mJy beam$^{-1}$ and 5 out of 21 sources (24\\%) are detected at a significance of $\\geq3\\sigma$. Stacking all sources leads to a statistical detection of $\\langle S_{1.2~\\rm mm} \\rangle = 0.96\\pm0.11$~mJy and stacking only the non-detections also yields a statistical detection, with $\\langle S_{1.2~\\rm mm} \\rangle = 0.51\\pm0.13$~mJy. At the typical redshift of the sample, $z=2$, 1~mJy corresponds to a far-infrared luminosity \\lfir$\\sim4\\times10^{12}$~\\lsol. If the far-infrared luminosity is powered entirely by star-formation, and not by AGN-heated dust, then the characteristic inferred star-formation rate is $\\sim$700~\\msolyr.  This far-infrared luminosity implies a dust mass of \\mdust$\\sim3\\times10^{8}$~\\msol, which is expected to be distributed on $\\sim$kpc scales. We estimate that such large dust masses on kpc scales can plausibly cause the obscuration of the quasars.  Combining our observations at 1.2~mm with mid- and far-infrared data, and additional observations for two objects at 350~\\mum\\, using SHARC-II, we present dust SEDs for our sample and derive a mean SED for our sample. This mean SED is not well fitted by clumpy torus models, unless additional extinction and far-infrared re-emission due to cool dust are included. This additional extinction can be consistently achieved by the mass of cool dust responsible for the far-infrared emission, provided the bulk of the dust is within a radius $\\sim$2-3~kpc.  Comparison of our sample to other samples of $z\\sim2$ quasars suggests that obscured quasars have, on average, higher far-infrared luminosities than unobscured quasars. There is a hint that the host galaxies of obscured quasars must have higher cool-dust masses and are therefore often found at an earlier evolutionary phase than those of unobscured quasars. For one source at $z=2.767$, we detect the CO(3-2) transition, with $S_{\\rm CO}\\Delta \\nu=$630$\\pm$50 mJy~\\kms, corresponding to $L_{\\rm CO (3-2)}=$3.2$\\times10^{7}$~\\lsol, or a brightness-temperature luminosity of $L'_{\\rm CO (3-2)}=2.4\\times10^{10}$ K~\\kms~pc$^{2}$. For another source at $z=4.17$, the lack of detection of the CO(4-3) line suggests the line to have a brightness-temperature luminosity $L'_{\\rm CO (4-3)}<1\\times10^{10}$ K~\\kms~pc$^{2}$. Under the assumption that in these objects the high-J transitions are thermalised, we can estimate the molecular gas contents to be $M_{\\rm H_{2}}=1.9\\times10^{10}$ \\msol and $<8\\times10^{9}$ \\msol, respectively.  The estimated gas depletion timescales are $\\tau_{\\rm g}=4$~Myr and $<$16~Myr, and low gas-to-dust mass ratios of \\mgas$/$\\mdust$=19$ and $\\leq8$ are inferred. These values are at the low end but consistent with those of other high-redshift galaxies. ",
        "introduction": "Quasars are believed to be powered by supermassive black holes (SMBHs) rapidly growing by accreting matter at a large fraction of their Eddington rate\\footnote{In this article we use the term quasar for the most powerful active galactic nuclei (AGNs). In the case of unobscured AGNs, the dividing line is taken at a B-band magnitude of $M_{\\rm B}=-23.5$ which, assuming the typical quasar SED of \\citet{1994ApJS...95....1E}, corresponds to $L_{\\rm bol}\\grtsim 2\\times10^{12}$~\\lsol. We will therefore also refer to obscured AGNs as obscured quasars if they have $L_{\\rm bol}\\grtsim 2\\times10^{12}$~\\lsol.}. The accreted matter forms a disk, is heated to temperatures of $\\sim10^{3}-10^{5}$~K, and emits thermal emission at optical and ultraviolet wavelengths. Clouds of gas ionised by this radiation will emit line radiation, and lines which are relatively close to the black hole have rapid orbits which will show broad velocity dispersions ($>$2000 \\kms), while clouds further away will show narrower lines.  Dust can survive when the equilibrium temperature with the ultraviolet photons is $\\lesssim$2000~K. This dust absorbs optical and ultraviolet radiation and re-emits it at infrared wavelengths characteristic of the dust temperature. If the line of sight of an observer to the central region is blocked by dust, the optical and ultraviolet continuum, as well as the broad lines, will not be observable. Such a source is known as an obscured quasar.  While radio-loud obscured quasars have been known for a long time, in the form of high-excitation narrow-line radio galaxies \\citep[see ][ for a review]{1995PASP..107..803U}, the population of radio-quiet obscured quasars had remained elusive, with only a few individual objects known. This has changed recently, when large numbers of these sources have been found in the spectroscopic database of the Sloan Digital Sky Survey \\citep{2003AJ....126.2125Z}. Samples of spectroscopically-confirmed obscured quasars have now also been identified in surveys using X-ray selection \\citep[e.g.][]{2004ApJS..155..271S,2005AJ....129..578B} as well as mid-infrared, sometimes combined with radio, selection \\citep[e.g.][]{2005ApJ...622L.105H,2005Natur.436..666M,2007ApJ...669L..61L}. It seems that the obscured quasars are at least as common as the the unobscured quasars \\citep[e.g.][]{2008AJ....136.2373R}, and probably outnumber these by a $\\sim$2:1 ratio \\citep[e.g.][]{2005Natur.436..666M,2008ApJ...674..676M,2007ApJ...669L..61L,2008ApJ...675..960P}.   There is a large body of evidence supporting obscuration of quasars as an orientation effect, where obscured and unobscured quasars are intrinsically identical sources, only dust around the accretion disk covers certain lines of sight so that the broad emission lines and the optical and ultraviolet continuum are not visible \\citep[the dusty torus of the unified schemes, e.g.][]{1995PASP..107..803U}. However, dust distributed within the host galaxy can also cause obscuration, particularly if the galaxy is at an early evolutionary phase and is rich in gas and dust \\citep[see e.g.][]{1982ApJ...256..410L,1988ApJ...325...74S,1999MNRAS.308L..39F}.   The importance of galaxy-scale dust is evident in recent studies of quasars where obscuration cannot be solely assigned to dust from a torus around the accretion disk: in some objects even the narrow-lines are obscured \\citep{2005Natur.436..666M,2006MNRAS.370.1479M,2006ApJ...645..115R}. The jets emanating from some obscured quasars are face-on \\citep{2006MNRAS.373L..80M,2007ApJ...667L..17S,2009arXiv0905.1605K}, something that is not expected for quasars obscured by the torus, but can be explained if the obscuring dust is on a larger ($\\sim$kpc) scale. The deep silicate absorption features and spectral energy distributions (SEDs) suggest foreground extinction is sometimes present \\citep{2007ApJ...654L..45L,2008ApJ...675..960P}. This all provides indirect evidence suggesting obscuring dust on $\\sim$~kpc scales, something which has been spectroscopically confirmed through the measurement of the H$\\alpha$ to H$\\beta$ Balmer decrement \\citep {2007ApJ...663..204B}.  Dust present on kpc scales is expected to be relatively cool ($T\\sim$50~K) and will emit thermally in the far-infrared ($\\lambda \\grtsim$40~\\mum), regardless of whether it is heated by young stars or by the central active galactic nucleus \\citep[AGN, see e.g.][ and Section~\\ref{sec:scale}]{1987ApJ...320..537B,1989ApJ...347...29S,2004A&A...421..129S}. Observing this dust around the peak of its emission (around rest-frame $\\sim$65~\\mum) is not optimal, since warmer dust can still contribute significantly there. The Rayleigh-Jeans tail of the emission, however, will have a much smaller contribution from warm dust, and has the additional advantage of being observable from the ground, for example in the atmospheric windows at 850~\\mum\\, or 1.2~mm.  In this paper we present continuum observations of a sample of 21 high-redshift ($z\\grtsim$2) radio-intermediate obscured quasars at 1.2~mm, combined with other infrared and submillimetre data, as well as a search for CO in two of the sources.  The sample was selected to approximately match the break in the unobscured luminosity function at $z\\sim2$ \\citep[$M_{\\rm B}\\sim$-25.7, ][]{2004MNRAS.349.1397C}, so in terms of radiation these quasars represent the energetically-dominant population around the peak of the quasar activity. The sources are, however, radio-intermediate in that their radio luminosities are slightly higher than those expected from radio-quiet quasars \\citep[with \\lrada\\, $\\sim10^{24}$~\\whzsr][]{2006MNRAS.373L..80M}. Nevertheless, observations with very large baseline interferometry suggest the relative importance of the compact cores makes them more similar to the genuinely radio-quiet population than to the radio galaxies and radio-loud quasars \\citep[see][]{2009arXiv0905.1605K}. Section~\\ref{sec:obs} summarises the observations and data reduction. In Section~\\ref{sec:res} we show the resulting fluxes, inferred luminosities, star-formation rates, dust masses, characteristic scale of the cool dust and discuss the possible obscuration by kpc-scale dust. Section~\\ref{sec:broad} shows the results of broad-band SED fitting and Section~\\ref{sec:mean} presents the mean SED for our sample. In Section~\\ref{sec:co} we discuss the molecular gas observations. We compare our sample to other millimetre or submillimetre observations of $z\\sim2$ quasars in Section~\\ref{sec:comp}. The discussion and summary are found in Section~\\ref{sec:disc}. ",
        "conclusions": "\\label{sec:disc}   We have observed a sample of $z\\grtsim2$ radio-intermediate obscured quasars at 1.2~mm (250~GHz) using MAMBO. The detection rate is 5 out of 21 (24\\%). Sources with narrow lines seem to have a higher detection rate than sources with no narrow lines, but with such small numbers the difference is not significant. Larger samples will be required to study any differences.  Stacking leads to a statistical detection of $\\langle S_{1.2~\\rm mm} \\rangle = 0.96\\pm0.11$. Even if only the non-detections are stacked, they still yield a statistical detection, with $\\langle S_{1.2~\\rm mm} \\rangle = 0.51\\pm0.13$. Thus, the typical flux density of this sample is $\\sim$0.5-1.0~mJy, and this corresponds to a far-infrared luminosity $\\sim4\\times10^{12}$~\\lsol.  If the far-infrared luminosity is powered entirely by star-formation, and not by AGN-heated dust, then the typical inferred star-formation rate is $\\sim$700~\\msolyr. This large star-formation rate is comparable to those inferred in ULIRGs and SMGs, the most powerful starburst galaxies known.  The observations at 1.2~mm also allow us to estimate the mass of the cool dust, and we find a typical mass of $\\sim3\\times10^{8}$~\\msol. If heated by young stars, this dust is expected to be distributed on $\\sim$2~kpc scales (the scale found for SMGs). If it is heated by the central AGN, with \\lbol$\\sim10^{13}$~\\lsol, then it is expected to be distributed on scales $\\lesssim$10~kpc.    Indeed, we estimate that such a large mass of cool dust is capable of causing alone the large values of \\av\\, in this sample, without help of an obscuring torus.  Combining our observations at mid-infrared and millimetre wavelengths, we present dust SEDs for our sample, and derive a typical SED for our sample of high-redshift obscured quasars. The clumpy tori models cannot reproduce the SED. However, clumpy tori models with an additional screen of cold dust (and the corresponding far-infrared emission) do reproduce this mean SED.  The amount of extinction from this additional screen is derived from the cool dust mass, from the observations at 1.2~mm. The depth of the silicate feature can be consistently achieved by the inferred cool dust mass provided that the dust is on scales $\\lesssim$2~kpc, which is in excellent agreement with the typical radius of the far-infrared emission in SMGs.  Again, this lends support to kpc-scale dust along the host galaxy playing an important role in the obscuration of these sources.  Obscuration by dust on kpc scales would also explain why in about half of the sample the narrow emission lines from the central AGN are not detectable, as well as why in some cases the jets seem to be pointing towards us, something unexpected if they are only obscured by the torus of the unified schemes.  However, we remind the reader that unobscured quasars have also been found to have similar cool-dust masses, and that in our sample there is not a one-to-one correspondence between the presence (or lack of) narrow lines and detection at 1.2~mm. The presence of a large mass of dust does not guarantee obscuration.  A clumpy dusty interstellar medium where different lines of sight have vastly different optical depths is a likely explanation for why some quasars with large dust masses are obscured while others are not.  When comparing to other samples of high-redshift quasars, we find that the obscured quasars probably have a higher fraction of their luminosity emerging at far-infrared wavelengths, compared to unobscured quasars. The far-infrared emission is isotropic, so that this difference cannot be ascribed to orientation-dependent obscuration: obscured quasars are likely to have higher far-infrared luminosities and cool-dust masses than unobscured quasars, suggesting the host galaxies are in a dustier, presumably earlier,  phase.   Additionally, we have searched for molecular gas in two sources at $z=2.767$ and 4.169, using the CO (3-2) and (4-3) transitions. In the $z=2.767$ source we detect a line with $L_{\\rm CO (3-2)}=$3.2$\\times10^{7}$~\\lsol\\, (equivalent to a brightness-temperature luminosity of $L'_{\\rm CO (3-2)}=2.4\\times10^{10}$ K~\\kms~pc$^{2}$). In the other source, the lack of detection suggests a line luminosity $L_{\\rm CO (4-3)}<$3$\\times10^{7}$~\\lsol\\, ($L'_{\\rm CO (4-3)}<1\\times10^{10}$ K~\\kms~pc$^{2})$. Under the assumption that in these objects the (3-2) and (4-3) transitions are thermalised, we can estimate the molecular gas contents to be $M_{\\rm H_{2}}=1.9\\times10^{10}$ and $<8\\times10^{9}$ \\msol, respectively.  The estimated gas depletion timescales are $\\tau_{\\rm g}=4$ and $<$16~Myr, and low gas-to-dust mass ratios of \\mgas$/$\\mdust$=19$ and $\\leq20$ are inferred. A dynamical mass of $M_{\\rm dyn} {\\rm sin^{2}}i=6\\times10^{9}$~\\msol\\, is estimated from the CO(3-2) detection."
    },
    "0910/0910.4733_arXiv.txt": {
        "abstract": "{ The solar abundances have undergone a major downward revision in the last decade, reputedly as a result of employing 3D hydrodynamical simulations to model the inhomogeneous structure of the solar photosphere. The very low oxygen abundance advocated by \\citet{asplund04}, A(O)=8.66, together with the downward revision of the carbon and nitrogen abundances, has created serious problems for solar models to explain the helioseismic measurements.  In an effort to contribute to the dispute we have re-derived photospheric abundances of several elements independently of previous analysis. We applied a state-of-the art 3D (CO5BOLD) hydrodynamical simulation of the solar granulation as well as different 1D model atmospheres for the line by line spectroscopic abundance determinations. The analysis is based on both standard disc-centre and disc-integrated spectral atlases; for oxygen we acquired in addition spectra at different heliocentric angles. The derived abundances are the result of equivalent width and/or line profile fitting of the available atomic lines. We discuss the different granulation effects on solar abundances and compare our results with previous investigations. According to our investigations hydrodynamical models are important in the solar abundance determination, but are not responsible for the recent downward revision in the literature of the solar metallicity. ",
        "introduction": "In this work we would like to face the most common questions we are confronted with in the analysis of the photospheric solar abundances: \\begin{itemize} \\item ``Are 3D models important in the abundances determination?'' \\item ``Are 3D models responsible for the downward revision of the solar metallicity?'' \\end{itemize} Our answer to the first question is yes. As we know from previous investigations \\citep{zolfito}, 3D solar metallicity models do not experience the over-cooling in the external layers, not detected in 1D models, that metal-poor 3D models show. One could then expect that 3D models are not fundamental for the solar abundance determinations. If for some elements (P, Eu, Hf) the granulation effects are in fact negligible, this is not the case for others, such as Fe, Th, and  also oxygen. On top of that one should not forget that 1D models require some input parameters (mixing-length parameter, \\mlp, and microturbulence, $\\xi_{\\rm micro}$) implicit in 3D models; so that in the case such parameters are fundamental (\\mlp\\ for C, O, and Fe, $\\xi_{\\rm micro}$ for N) 3D models should be highly preferred.   \\begin{table} \\caption{Influence of the microturbulence on the 3D corrections.} \\label{microt} { \\begin{tabular}{ccc} \\hline \\noalign{\\smallskip} $\\xi _{\\rm \\xx}$ & \\multicolumn{2}{c}{$\\left({\\rm A(Y)_{\\rm 3D}}-{\\rm A(Y)_{\\rm \\xx}}\\right)$ [dex]} \\\\ \\kms\\ & Flux & Intensity \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} 0.6 & +0.036 &  - -\\\\ 0.9 &  - -   & +0.037\\\\ 1.2 & +0.130 & - - \\\\ 1.5 &  - -   & +0.139\\\\ \\noalign{\\smallskip} \\hline \\end{tabular} } \\end{table}   \\begin{table*} \\caption{Photospheric solar abundances.} \\label{sunabbo} \\begin{center} {% \\begin{tabular}{lrlllll} \\hline \\noalign{\\smallskip} EL & N  & \\cobold\\ & AG89 & GS98 & AGS05 & AGSS09\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} Li & 1  & $1.03\\pm 0.03$ & $1.16\\pm 0.10$ & $1.10\\pm 0.10$        & $1.05\\pm 0.10$        & $1.05\\pm 0.10$ \\\\ C  & 43 & $8.50\\pm 0.06$ & $8.56\\pm 0.04$ & $8.52\\pm 0.06$        & $8.39\\pm 0.05$        & $8.43\\pm 0.05$ \\\\ N  & 12 & $7.86\\pm 0.12$ & $8.05\\pm 0.04$ & $7.92\\pm 0.06$        & $7.78\\pm 0.06$        & $7.83\\pm 0.05$ \\\\ O  & 10 & $8.76\\pm 0.07$ & $8.93\\pm 0.035$& $8.83\\pm 0.06$        & $8.66\\pm 0.05$        & $8.69\\pm 0.05$ \\\\ P  & 5  & $5.46\\pm 0.04$ & $5.45\\pm 0.04$ & $5.45\\pm 0.04$        & $5.36\\pm 0.04$        & $5.41\\pm 0.03$ \\\\ S  & 9  & $7.16\\pm 0.05$ & $7.21\\pm 0.06$ & $7.33\\pm 0.11$        & $7.14\\pm 0.05$        & $7.12\\pm 0.03$ \\\\ K  & 6  & $5.11\\pm 0.09$ & $5.12\\pm 0.13$ & $5.12\\pm 0.13$        & $5.08\\pm 0.07$        & $5.03\\pm 0.09$ \\\\ Fe & 15 & $7.52\\pm 0.06$ & $7.67\\pm 0.03$ & $7.50\\pm 0.05$        & $7.45\\pm 0.05$        & $7.50\\pm 0.04$ \\\\ Eu & 5  & $0.52\\pm 0.03$ & $0.51\\pm 0.08$ & $0.51\\pm 0.08$        & $0.52\\pm 0.06$        & $0.52\\pm 0.04$ \\\\ Hf & 4  & $0.87\\pm 0.04$ & $0.88\\pm 0.08$ & $0.88\\pm 0.08$        & $0.88\\pm 0.08$        & $0.85\\pm 0.04$ \\\\ Os & 3  & $1.36\\pm 0.19$ & $1.45\\pm 0.10$ & $1.45\\pm 0.10$        & $1.45\\pm 0.10$        & $1.40\\pm 0.08$ \\\\ Th & 1  & $0.08\\pm 0.03$ & $0.12\\pm 0.06$ & ${\\it 0.09\\pm 0.02}$  & ${\\it 0.06\\pm 0.05}$  & $0.02\\pm 0.10$ \\\\ &    &                &                &                       &                       & \\\\ Z  &    & 0.0153         & 0.0189         & 0.0171                & 0.0122                & 0.0134 \\\\ Z/X&    & 0.0209         & 0.0267         & 0.0234                & 0.0165                & 0.0183 \\\\  \\noalign{\\smallskip} \\hline \\end{tabular} } \\end{center} Note: AG89=\\citet{anders89}, GS98=\\citet{grevesse98}, AGS05=\\citet{sunabboasp}, and AGSS09=\\citet{asplund09}. \\end{table*}     Our answer to the second question is no. We would like to remark that this answer depends on the 1D reference model one is considering. For all the elements, except N and Th, our 3D model gives an abundance larger than the 1D model. Th is an exception because its lower 3D abundance is related to the fact that the line is on the red wing of a stronger Fe-Ni blend, and the asymmetry of this blend, correctly taken into account in 3D analysis, lowers the Th contribution \\citep{thhf}. For N all the lines considered are of high excitation energy, so their 1D abundance is very sensitive on the choice of \\mlp. To try to catch the effects of granulation on the abundances determination, we selected 1D models sharing the micro-physics with our 3D solar model atmosphere. But the comparison of abundances derived from 1D and 3D models is not unambiguously defined, relying on the choice of \\mlp\\ and  $\\xi_{\\rm micro}$. ",
        "conclusions": "In the light of our work we think that the use of 3D models in the solar abundance determination is useful. This is for the following reasons: \\begin{itemize} \\item no constraint is necessary on \\mlp\\ and $\\xi_{\\rm micro}$; \\item a difference of few hundreds of dex in the abundance, negligible for the majority of the stellar analysis, is important in the solar context. \\end{itemize}    \\balance"
    },
    "0910/0910.1320_arXiv.txt": {
        "abstract": "We report a relation between radio emission in the inner jet of the Seyfert galaxy \\objectname{3C~120} and optical continuum emission in this galaxy.  Combining the optical variability data with multi-epoch high-resolution very long baseline interferometry observations reveals that an optical flare rises when a superluminal component emerges into the jet and its maxima is related to the passage of such component through the location a stationary feature at a distance of $\\approx$1.3\\,parsecs from the jet origin. This indicates that a significant fraction of the optical continuum produced in \\objectname{3C~120} is non-thermal and it can ionize material in a sub-relativistic wind or outflow. We discuss implications of this finding for the ionization and structure of the broad emission line region, as well as for the use of broad emission lines for determining black hole masses in radio-loud AGN. ",
        "introduction": "In the current astrophysical paradigm for active galactic nuclei (AGN), each constituent of an AGN contributes to a specific domain in the spectral energy distribution (SED).  The SED of some AGN, show a significant amount of energy excess at UV/optical wavelengths, which is commonly attributed to be produced in an accretion disk around the putative central supermassive black hole (BH). Surrounding the accretion disk at $\\lesssim 1$ pc, there is a central broad line region (BLR) formed by high density gaseous clouds orbiting the central BH. The thermal UV/optical emission radiated from the disk is thought to be the prime source of variable optical continuum and a dominant factor for the ionization of the BLR material. In radio-loud AGN however, the broad-band continuum can also have a substantial contribution from non-thermal synchrotron radiation generated in the relativistic jet (D'Arcangelo et al. 2007, Marscher et al. 2008, Soldi et al. 2008). To understand the mechanism and properties of the broad-line generation in such objects, it is pivotal to be able to localize and identify the region where the bulk of the non-thermal optical continuum emission is produced.  An efficient way to identify a region responsible for the production of the variable optical continuum emission in radio-loud AGN is to combine long-term optical monitoring and regular very long baseline interferometry (VLBI) observations. A link between the optical emission and relativistic jet in 3C~120 was suggested by Belokon (1987), based on a correlation found between the optical outburst and ejections of plasma clouds (jet components) in the jet. Using the radio and optical data available from 14 years monitoring of the radio galaxy 3C\\,390.3, Arshakian et al. (2008, 2009) found that ejections of new components can be associated with optical flares occurring on time scales from months to years and reaching their maxima during passages of the jet components through a stationary emitting region in the inner jet, at a distance of $\\sim 0.5$\\,pc from the jet origin. This indicates that the variable optical continuum emission can be generated in the innermost part of the jet, in the region located upstream from the stationary feature. This has an important implication for the existence of a non-virialized outflowing broad-line region associated with the jet.   \\begin{figure*}[t!] \\includegraphics[width=\\textwidth]{f1.eps} \\caption{ Compact jet in 3C~120 observed on January 7, 2005, with the VLBA at 15~GHz (2cm). Shaded ellipse in the lower left corner marks the restoring point-spread function (beam) with the FWHM of 0.53 $\\times$ 1.26 mas oriented at an angle of three degrees (clockwise rotation). The peak flux density in the image is 974 mJy/beam and the rms noise is 0.26 mJy/beam. Contours are drawn at 1, $\\sqrt{2}$, 2...\\% of the lowest contour shown at 0.9 mJy/beam. The jet structure is parametrized by a set of two-dimensional, circular Gaussian features obtained from fitting the visibility amplitudes and phases (Pearson 1997). The model-fit components (C4-C12) located within the inner 14 mas from the core are shown in Figure \\ref{Fig:3c120_kinematic_model}.  Label marks show the location of the model-fit components and the two stationary components D and S1 are highlighted by open circles.} \\label{Fig:3c120_jet} \\end{figure*}    It should be noted that stationary and low-pattern speed features have been recently reported in parsec-scale jets of a number of prominent radio sources ({\\em e.g.}, Kellermann et al. 2004, Savolainen et al. 2006, Lister et al. 2009b), alongside faster, superluminally moving components embedded in the same flows (Kellermann et al. 2004, Jorstad et al. 2005). Although specific geometric conditions and extremely small viewing angles may lead to formation of such features in relativistic flows (Alberdi et al. 2000), it is more likely that these features represent standing shocks (for instance, recollimation shocks in an initially over-pressurized outflow; G\\'omez et al. 1995, Perucho \\& Mart\\'i 2007). Such standing shocks may play a major role in accelerating particles near the base of the jet (Mandal \\& Chakrabarti 2008; Becker et al. 2008), and could be responsible for the persistent high levels of polarization in blazars (D'Arcangelo et al. 2007; Marscher et al. 2008).  It is important to understand whether the correlations found in Arshakian et al. (2008, 2009) are characteristic only for the 3C\\,390.3 or common for all radio galaxies and quasars. In this letter, we combined archival VLBI monitoring data and data available from photometric monitoring of 3C\\,120 to study the link between properties of the superluminal jet and optical continuum flares.  Throughout the paper, a flat $\\Lambda$CDM cosmology is assumed, with the Hubble constant H$_{0}= 70$ km s$^{-1}$ Mpc$^{-1}$ and matter density $\\Omega_{m}=0.3$. This corresponds to a linear scale of 0.658\\,pc/mas, at the redshift $z=0.033$ of \\objectname{3C~120}. ",
        "conclusions": "To investigate the link between optical continuum variability and subparsec-scale jet in the radio-loud galaxy 3C\\,120, we combined the radio VLBI (15\\,GHz) and optical photometry (B-band) data observed during the period of nearly eight years. We found a significant correlation between the jet kinematics on parsec-scales and optical flares on scales from several months to about two years. All optical flares are associated with the jet ejection events: the flare rises after the epoch of ejection of a new jet component (at D) and it reaches the maximum around the epoch at which the ejected radio knot passes the stationary radio component (S1) downstream the jet. The radio passages of new components through S1 delay the maxima of optical flares on the average by $\\la 0.1$ yr. These results confirm correlations found in Arshakian~et~al.~(2008,~2009) for the radio galaxy 3C\\,390.3 and support the idea that the correlation between optical flares and kinematics of the jet could be a common feature for all radio-loud galaxies and quasars.  The link between optical continuum variability and kinematics of the parsec-scale jet is interpreted in terms of optical flares generated by disturbances in the jet flow while moving from the stationary component D of the jet to the stationary component S1 located at a distance of about one parsec. Modeling of optical flares as synchrotron flares caused by density variations in the relativistic flow showed that all optical flares require at most a factor of seven increase of the particle density. The variation of the particle density at subparsec-scales should be smaller if consider a more realistic model in which the magnetic field and/or Lorentz factor of the jet increase downstream of the jet, in the acceleration/collimation zone, between D and S1. In this scenario, the correlated X-ray and optical continuum emission are produced in the relativistic jet at some distance ($< 1$ pc) from the central black hole rather than in the accretion disk. \\\\  Establishing such a relation may have strong implications for the physics of the central regions in radio-loud AGN and AGN in general (see also Arshakian et al. 2009). The link between the optical continuum and radio jet challenges the existing models in which the optical continuum and broad-line emission are both localized around the disk or near the central black hole of an AGN. It also questions the common assumptions about the central engine used, in particular, for the reverberation mapping technique (Peterson et al. 2002), which combines the broad emission line width with the size of the line-emitting region estimated from the optical continuum luminosity. In 3C\\,120, a substantial fraction of the ionizing continuum is produced in the relativistic jet and it illuminates thermal material most likely organized in a sub-relativistic outflow. The time delays and profile widths measured in radio-loud AGN may reflect not only the Keplerian motion of the emitting line gas, but also an outflowing component of the gas accelerated by non-gravitational forces. This can lead to large errors in estimates of black hole masses made from monitoring of the broad emission lines.  In order to address these issues, we are carrying out a long-term spectropolarimetry campaign for the \\objectname{3C~120} and several other nearby radio galaxies aimed to determine the amount of non-thermal optical continuum and to tackle evidence for non-virial motions in the BLR. Follow-up high resolution radio observations at other frequencies and epochs, can help to characterize the nature of the components D and S1.    We thank anonymous referees for a number of useful comments and suggestions which significantly improved this manuscript.  This work was supported by CONACYT research grant 54480 (Mexico) and the Russian Foundation for Basic Research (Grant No. 09-02-01136a) .  JLT acknowledges support from the CONACYT program for PhD studies, the International Max-Planck Research School for Radio and Infrared Astronomy at the Universities of Bonn and Cologne and the Deutscher Akademischer Austausch Dienst (DAAD) for a short-term scholarship in Germany. This research has made use of data from the MOJAVE database that is maintained by the MOJAVE team (Lister et al. 2009a). The VLBA is an instrument of the National Radio Astronomy Observatory, a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc."
    },
    "0910/0910.1927_arXiv.txt": {
        "abstract": "Although they are only minor constituents of the interstellar medium, halogen-containing molecules are of special interest because of their unique thermochemistry.  Here, we present a theoretical study of the chemistry of interstellar molecules containing the halogen elements chlorine and fluorine.  We have modeled both diffuse and dense molecular clouds, making use of updated estimates for the rates of several key chemical processes.  We present predictions for the abundances of the three halogen molecules that have been detected to date in the interstellar medium:  HF, CF$^+$ and HCl.  As in our previous study of fluorine-bearing interstellar molecules, we predict HF to be the dominant gas-phase reservoir of fluorine within both diffuse and dense molecular clouds; we expect the {\\it Herschel Space Observatory} to detect widespread absorption in the HF $J=1-0$ transition.  Our updated model now overpredicts the CF$^+$ abundance by a factor $\\simgt 10$ relative to observations of the Orion Bar; this discrepancy has widened because we now adopt a laboratory measurement of the CF$^+$ dissociative recombination rate that is smaller than the estimate we adopted previously.  This disagreement suggests that the reaction of C$^+$ with HF proceeds more slowly than the capture rate assumed in our model; a laboratory measurement of this reaction rate would be very desirable.  Our model predicts diffuse cloud HCl abundances that are similar to those predicted previously and detected tentatively toward $\\zeta$~Oph. Two additional species are potentially detectable from photodissociation regions: the  $\\rm H_2Cl^+$ and $\\rm HCl^+$ molecular ions.  Ortho-$\\rm H_2Cl^+$ has its lowest-lying transition in the millimeter spectral region observable from the ground, and the lowest rotational transition of HCl$^+$ is observable with {\\it Herschel}'s HIFI instrument. ",
        "introduction": "Of the $\\sim$ 150 distinct molecules detected thus far in interstellar clouds and circumstellar outflows, over  90$\\%$ contain only hydrogen and/or elements in groups IVA, VA, and VIA\\footnote{Here we use the traditional designations used in the United States.  The current IUPAC designations for groups IVA, VA, and VIA are respectively 14, 15, and 16; those for groups IA, IIA, IIIA and VIIA are 1, 2, 13, and 17} of the periodic table: viz.\\ carbon, nitrogen, oxygen, silicon, phosphorus, and sulfur.  Of the remaining 12 detected molecules, eight contain elements in groups IA, IIA and IIIA (viz.\\ sodium, potassium, magnesium and aluminum) and are found exclusively in circumstellar envelopes (e.g.\\ Ziurys 1996, and references therein) with abundances that are set in cool stellar photospheres; one contains iron (FeO, detected toward the Sgr B2 interstellar gas cloud by Walmsley et al.\\ 2002); and the remaining 3 are interstellar molecules containing elements of the halogen group (VIIA): HCl, HF, and CF$^+$.  Although they are only minor constituents of the interstellar medium, halogen-containing molecules are of special interest because of their unique thermochemistry.  The basic behavior of any atom in the interstellar medium is determined by two or three key thermochemical considerations.  First, the first ionization potential (IP) determines whether the atom will be neutral or ionized within diffuse molecular clouds: an atom X with IP(X) $>$ IP(H) = 13.6 eV will be shielded by atomic hydrogen from the interstellar radiation field (ISRF) and will be predominantly neutral, while an atom with IP(X) $<$ IP(H) will be predominantly ionized.  Second, the dissociation energy of the neutral hydride, $D_0$(HX), determines whether X can react exothermically with H$_2$: atoms for which $D_0$(HX) $>$ $D_0$(H$_2$) = 4.48 eV can react exothermically with H$_2$ to form HX, while atoms with  $D_0$(HX) $ < D_0$(H$_2$) cannot.  Third (and of importance mainly for atoms with IP $<$ 13.6 eV),  the dissociation energy of the hydride molecular ion, $D_0$(HX$^+$), similarly determines whether the ion X$^+$ can react exothermically with H$_2$ to form HX$^+$.  Fluorine and chlorine\\footnote{In this paper we combine our attention to Cl and F; bromine and iodine have much smaller cosmic abundances than chlorine and fluorine, and molecules containing bromine and iodine are unlikely to be detectable with current technology.} are unique in regard to the second and third criteria, each in a different way: F is the only atom -- and Cl$^+$ is the only ion with an appearance potential $<$ 13.6 eV -- that can react exothermically with H$_2$ to form a diatomic hydride.  The uniqueness of these elements are presented graphically in Figure 1, which shows $D_0$(HX) for each atom X in the first three rows of the periodic table.  The molecular binding energy of HX is largest for fluorine and decreases, nearly monotonically, as one moves leftwards and downwards in the periodic table.  In Figure 2, $D_0$(HX) and $D_0$(HX$^+$) are shown for all the elements of which interstellar molecules have been detected.  Red squares apply to atoms with  IP $<$ IP(H) and blue squares to those with IP $>$ IP(H).  With the exception of the halogen elements, all the elements fall into two classes:  (1) {\\it Oxygen and nitrogen, with IP $>$ IP(H), $D_0$(HX) $<$ $D_0$(H$_2$) and $D_0$(HX$^+$) $>$ $D_0$(H$_2$).} These elements are predominantly neutral in the diffuse ISM.  Their atoms do not react with H$_2$ except at elevated temperatures that can drive endothermic reactions, and their chemistry is driven by cosmic ray ionization.  (2) {\\it Carbon, silicon, sulfur, and phosphorus, with IP $<$ IP(H), $D_0$(HX) $<$ $D_0$(H$_2$) and $D_0$(HX$^+$) $<$ $D_0$(H$_2$).}  These elements are predominantly ionized in the diffuse ISM.  Their ions do not react with H$_2$ except at elevated temperatures that can drive endothermic reactions.  Chlorine and fluorine form two additional classes with one member each:  (3) {\\it Fluorine, with IP $>$ IP(H) and $D_0$(HX) $>$ $D_0$(H$_2$).}  Fluorine is predominantly neutral in the diffuse ISM, and can react exothermically with H$_2$ to form HF.  (4) {\\it Chlorine, with IP $<$ IP(H), $D_0$(HX) $<$ $D_0$(H$_2$) and $D_0$(HX$^+$) $>$ $D_0$(H$_2$).}  Chlorine is predominantly ionized in atomic clouds.  Cl$^+$ can react exothermically with H$_2$ to form HCl$^+$.  The chemistry of interstellar fluorine has been considered in a previous study by Neufeld, Wolfire \\& Schilke (2005; hereafter NWS05), while that of chlorine has been discussed in several studies, most recently those of Schilke, Wang \\& Phillips (1995) and Amin (1996).  In the present paper, we extend and update these previous investigations to account for recent developments in experimental and theoretical studies of relevant chemical processes.  Our study is also motivated by the prospects for upcoming observations of halogen-containing molecules, particularly with the {\\it Herschel Space Observatory}. ",
        "conclusions": "As expected, the results shown in Figures 3 -- 7 indicate qualitatively-different behavior for the F- and Cl-bearing species.  The results for F-bearing molecules are qualitatively the same as those obtained previously by NWS05: once H$_2$ becomes abundant, HF becomes the dominant reservoir of gas-phase fluorine, with CF$^+$ the only other significant molecule containing fluorine.  Because of the lower dissociative recombination rate assumed for CF$^+$ in the present work, the predicted CF$^+$ abundances are somewhat larger.  The chemistry for Cl-bearing molecules is somewhat more complex, by contrast, and the computer-generated network diagrams shown in Figure 9, together with Figures 10 - 13 in the online materials, help elucidate the key pathways.  \\subsection{Identification of key chemical pathways}  In Figures 9 -- 13, circles represent the key Cl-bearing species, with a color scale indicating their relative abundance.  The color map extends from cyan to blue to purple to red (highest value), with grey circles denoting species that account for less than 10$^{-8}$ of the gas-phase chlorine abundance. Arrows indicate the key reactions linking the various species: again, the colors of the arrows indicate the reaction rate per unit volume. The results apply to a one-sided model with $n_{\\rm H} = 10^4\\,\\rm cm^{-3}$ and $\\chi_{UV} = 10^4$, and the different figures show the behavior at various depths into the cloud.  The depths were chosen to illustrate several different regimes.  1) $A_V = 0$:  At the cloud surface, there is no shielding and the abundance of molecules is small.  The Cl$^+$/Cl ratio is $\\sim 10^5$, being determined by the balance between photoionization and recombination. For the latter process, radiative recombination and charge transfer recombination with H -- although endothermic by $\\sim$~0.6 eV -- are of roughly equal importance.  2) $A_V= 0.8$:  Here, H$_2$ self-shielding has greatly increased the H$_2$ abundance.  Reaction with H$_2$ is now the dominant destruction mechanism for Cl$^+$, resulting in the formation of HCl$^+$.  However, the e/H$_2$ abundance ratio is still sufficient for dissociative recombination to dominate the destruction of HCl$^+$, although the competing reaction with H$_2$ produces a small amount of H$_2$Cl$^+$.  The dissociative recombination of HCl$^+$ leads to atomic Cl, which becomes the dominant gas-phase reservoir of chlorine. Beyond $A_V \\sim 0.8$, the abundances of the HCl$^+$ and H$_2$Cl$^+$ ions start to fall, along with that of $\\rm Cl^+$.  3) $A_V = 2.0$: The e/H$_2$ abundance ratio has now dropped sufficiently that HCl$^+$ reacts primarily with H$_2$, rather than electrons, forming H$_2$Cl$^+$.  HCl is produced by dissociative recombination of H$_2$Cl$^+$, although the dominant formation process for HCl is the slightly endothermic reaction of H$_2$ and Cl; this reaction is still possible at the $\\sim 200$~K temperature at this point within the PDR.  At $A_V = 2.0$, Cl and HCl account respectively for $\\sim 0.1 \\%$ and $\\sim 99.9 \\%$ of elemental chlorine in the gas phase.  4) $A_V = 4.0$: The gas temperature drops below the level at which HCl can be formed by reaction of H$_2$ with Cl, and the HCl abundance is lower than it was at $A_V = 2.0$.  The Cl$^+$/Cl ratio is now only $10^{-8}$, the photoionization rate of Cl having dropped dramatically.  Thus the formation of HCl$^+$ is dominated by reaction of H$_3^+$ with Cl, not H$_2$ with Cl$^+$.  The H$_3^+$ molecular ion is produced by the effects of cosmic rays; these ionize H$_2$, forming H$_2^+$, which then transfers a proton to H$_2$ to form H$_3^+$.  5) $A_V = 10.0$:  As at $A_V=4.0$, HCl$^+$ is produced primarily by reaction of H$_3^+$ with Cl, and reaction of  HCl$^+$ with H$_2$ then forms H$_2$Cl$^+$.  At this depth, the electron fraction is so low that dissociative recombination of H$_2$Cl$^+$ is no longer dominant.  H$_2$Cl$^+$ transfers a proton to a neutral species of higher proton affinity than HCl (e.g. CO or H$_2$O; this process is denoted by the arrow labeled ''N''), yielding HCl.  At this point, the external UV is almost completely attenuated, and the HCl destruction rate is very small.  HCl now accounts for several $\\times 10\\%$ of the gas-phase chlorine abundance, with atomic chlorine accounting for essentially all of the remainder.  Our results in this regime are in good agreement with those obtained previously by Schilke et al\\ (1995). H$_2$Cl$^+$ and CCl$^+$ are also present, but their abundances are just a few $\\times 0.01\\%$ of gas-phase elemental chlorine.  HCl reacts with H$_3^+$ to form H$_2$Cl$^+$, although this rapidly leads to reformation of HCl by proton transfer to other neutral species.  Several other destruction processes are of comparable importance for HCl destruction:  cosmic ray induced photodissociation, and reaction with the positive ions C$^+$ and He$^+$.  In this regime, the rates of HCl formation and destruction all scale linearly with the cosmic-ray ionization rate, with the result that the HCl abundance is almost independent of the latter.  Although the diagrams shown in Figure 9 -- 13 apply to just a single model, similar regimes exist for other values of $\\chi_{UV}$ and $n_{\\rm H}$.  Although the relevant regions move in or out (to larger or smaller $A_V$) as the assumed $\\chi_{UV} / n_{\\rm H}$ ratio is increased or decreased, the same qualitative behavior is observed.  In examining similar diagrams for the entire set of models, we have not identified any important reaction pathway that is not apparent in Figures 9 -- 13.  \\subsection{Observational implications}  To date, HF, CF$^+$, and HCl are the only halogen-bearing molecules to have been detected in the interstellar medium.  In regard to HF, our results are identical to those presented in NWS05: we predict HF to be the dominant gas-phase reservoir of fluorine within both diffuse and dense molecular clouds; we expect the {\\it Herschel Space Observatory} to detect widespread absorption in the HF $J=1-0$ transition. However, the abundances we predict for CF$^+$ lie a factor $\\sim 3$ above the prediction of NWS05, a direct consequence of the smaller dissociation rate adopted for CF$^+$ in the present study.  As noted by Neufeld et al. (2006), the CF$^+$ column densities inferred from observations of the Orion Bar were already a factor of $\\sim 4$ below the predictions of NWS05, and the discrepancy between theory and observation is now increased to more than an order of magnitude.  This disagreement may indicate that we have overestimated the rate coefficient for reaction of C$^+$ and HF. We are not aware of any laboratory studies of this reaction; the rate we adopt is simply the capture rate, and thus an upper limit on the true reaction rate.  Laboratory measurements of this key reaction would be very desirable.  HCl has been observed in both diffuse and dense molecular clouds.  In diffuse clouds, ultraviolet absorption studies have led to upper limits, and in one case, a tentative detection toward $\\zeta$ Oph (Federman et al.\\ 1995).  The latter result, obtained using the Goddard High Resolution Spectrograph on HST, yielded an HCl column density of $2.7 \\pm 1.0 (1\\sigma) \\times 10^{11} \\rm cm^{-2}$.  Given an atomic chlorine column density of $3.0 \\pm 1.0 (1\\sigma) \\times 10^{14} \\rm cm^{-2}$ for this sight-line (Federman et al.\\ 1995), the corresponding $N({\\rm HCl})/N({\\rm Cl})$ ratio is $9 \\times 10^{-4}$ (uncertain by a factor $\\sim 2$), a value in excellent agreement with predictions of the diffuse cloud models presented by vDB86.   Very similar results are obtained from our updated treatment of Cl chemistry.  In Figure 14, we show predictions from a series of models with $\\chi_{UV}$ ranging from $10^{-1}$ to 10$^2$.  All models apply to a slab with parameters appropriate to $\\zeta$~Oph, viz.\\ density $n_H = 10^{2.5} \\rm \\, cm^{-3}$ and total visual extinction of 0.8 mag.  The horizontal axis shows the H$_2$ column density, for which the measured value is $4.2 \\times 10^{20}$ along this sight-line, while the vertical axis shows $N({\\rm HCl})/N({\\rm Cl})$ (red curve), $N({\\rm Cl^+})/N({\\rm Cl})$ (orange) and $N({\\rm HCl^+})/N({\\rm Cl})$ (blue).  Squares located along the curves show the results for $\\chi_{UV} = 10^{-1}$ to 10$^2$ from right to left.  The H$_2$ column density along the $\\zeta$~Oph sight-line requires a $\\chi_{UV} \\sim 10^{0.5}$, in agreement with previous studies, and the predicted $N({\\rm HCl})/N({\\rm Cl})$ is in good agreement with the observations.  The almost exact agreement between the present model and that of vDB86 in regard to the predicted HCl abundance appears to result from the fortuitious cancellation of two changes in the photochemical network: the photodissociation rate for HCl and the photoionization rate for Cl are both increased by a factor $\\sim 2$. As in the vDB86 models, the predicted Cl$^+$ column density lies substantially below the observations.   Federman et al.\\ (1995) attributed this discrepancy to the presence of a significant Cl$^+$ column density within  HII regions along the sight-line. One important caveat must be noted.  The results shown in Figure 14 were obtained by assuming a branching ratio of only 10$\\%$ for the production of HCl following dissociative recombination of H$_2$Cl$^+$.  This value was initially invoked primarily to {\\it explain} the relative low HCl abundance derived from upper limits obtained with the {\\it Copernicus} satellite.  Models that we obtained for other branching ratios (0, 0.3 and 1.0) indicate that the HCl column density is linearly proportional to the branching ratio for values greater than 0.1  With a branching ratio of zero (i.e.\\ with no production of HCl from dissociative recombination of H$_2$Cl$^+$), a non-zero but negligible HCl column density ($N({\\rm HCl})/N({\\rm Cl}) \\sim 10^{-6}$) results from the endothermic reaction of H$_2$ with Cl.  In dense clouds, HCl has been unequivocally detected in several sources by means of submillimeter observations of the $J=1-0$ emission line.  First detected using the Kuiper Airborne Observatory toward Orion (Blake, Keene \\& Phillips 1985) and then Sgr B2 (Zmuidzinas et al.\\ 1995), the $J=1-0$ transition has also been observed using the ground-based Caltech Submillimeter Observatory (CSO) under good atmospheric conditions.  CSO observations have led to additional detections toward Orion A, Mon R2 (Salez, Frerking \\& Langer 1996) and several other sources (Phillips et al.\\ 2009), as well as mapping observations toward OMC-1 (Schilke, Phillips \\& Wang 1995).  Typical column densities derived for these dense clouds lie in the few $\\rm \\times 10^{13}$ to few $\\times 10^{14} \\, \\rm cm^{-2}$ range, in reasonable agreement with the predictions of our model for regions at high density exposed to strong UV radiation (Fig.\\ 5).  As noted in the observational studies cited above, depletion plays an important role in limiting the fractional abundance of HCl.  Our study identifies two additional Cl-bearing species that are potentially detectable: the molecular ions HCl$^+$ and H$_2$Cl$^+$.   These ions are isoelectronic with OH and H$_2$S respectively.  They are most abundant near cloud surfaces, where the photoionization rate for HCl is highest, and show column densities that are an increasing function of $\\chi_{UV} / n_{\\rm H}$.  Thus, PDRs subject to strong UV irradiation present attractive targets for searches for HCl$^+$ and H$_2$Cl$^+$ at millimeter and submillimeter wavelengths.  For example, in a PDR with $n_{\\rm H} = 10^4\\,\\rm cm^{-3}$ and $\\chi_{UV} = 10^4$, column densities of $5.7$ and $2.6\\times 10^{11} \\, \\rm cm^{-2}$ are predicted for HCl$^+$ and H$_2$Cl$^+$.  These values apply to sight-lines perpendicular to the illuminated surface; in edge-on PDRs, they can be significantly enhanced by limb-brightening.  The rotational spectrum of HCl$^+$ has been measured using laser magnetic spectroscopy (Lubic et al.\\ 1989).  The lowest-lying rotational transition, the $^2\\Pi_{3/2} J = 5/2 \\rightarrow 3/2$ multiplet near 207.6$\\mu$m, lies in a wavelength range accessible to the HIFI instrument on the {\\it Herschel Space Observatory}.  The rotational spectrum of H$_2$Cl$^+$ has also been the subject of a laboratory study at high spectral resolution; here, the measurements of Araki et al.\\ (2001) indicate that the lowest-lying rotational line of ortho-H$_2$Cl$^+$, the $1_{10}-1_{01}$ transition, lie near 189.2 and 188.4 GHz respectively for the  $\\rm H_2^{35}Cl^+$ and $\\rm H_2^{37}Cl^+$ isotopologues.  (These transitions are split into 6 and 4 hyperfine components, covering $\\sim 56$ and $14$~MHz respectively.)  This spectral region lies in the wing of the strong 183 GHz telluric water absorption line, but is nevertheless observable from ground-based observatories under favorable -- but by no means unusual -- atmospheric conditions.   In one of the  earliest studies of interstellar chlorine chemistry, the possibility of detecting HCl$^+$ in {\\it diffuse} molecular clouds was discussed by Jura (1974), who considered the $A^2\\Sigma^+ \\leftarrow X ^2\\Pi$ band that is accessible in the near-ultraviolet region.   Given an oscillator strength of $4.15 \\times 10^{-4}$ for the strongest vibrational band (Pradhan, Kirby \\& Dalgarno 1991), and given the abundances predicted in Figure 17, our results for $\\zeta$~Oph would imply equivalent widths of only $\\sim 4 \\times 10^{-6}\\,\\AA$, well below the limit of detectability.  Finally, we note that, unlike CF$^+$, the CCl$^+$ ion is typically predicted to have an extremely low abundance.  Although CCl$^+$ is produced rapidly by reaction of HCl and C$^+$, these two species show very little overlap: the HCl abundance is small in the surface layers where C$^+$ is abundant, unless the temperature is increased by shock heating, and the C$^+$ abundance is small in the deep cloud interiors where HCl is relatively abundant.  Enhanced CCl$^+$ abundances are possible when shock heating is present along with UV irradiation."
    },
    "0910/0910.3922_arXiv.txt": {
        "abstract": "Microlensing and occultation are generally studied in the geometric optics limit.  However, diffraction may be important when recently discovered Kuiper-Belt objects (KBOs) occult distant stars.  In particular the effects of diffraction become more important as the wavelength of the observation and the distance to the KBO increase. For sufficiently distant and massive KBOs or Oort cloud objects not only is diffraction important but so is gravitational lensing.  For an object similar to Eris but located in the Oort cloud, the signature of gravitational lensing would be detected easily during an occultation and would give constraints on the mass and radius of the object. ",
        "introduction": "\\citet{Bailey:1976p1677} first argued that small bodies in the distant solar system could be detected through stellar occultations. \\citet{1987AJ.....93.1549R} developed the treatment of occultation by irregular bodies including diffraction.  More recently with the discovery of the population of Kuiper Belt objects \\citep{Kuiper:1951p1773,Jewitt:1999p1844}, several research groups have begun searching for more distant and smaller bodies through occultations \\citep{Roques:2000p1661,2008AJ....135.1039B,Roques:2009p1666}. Because Kuiper Belt objects (KBOs) typically subtend small angles, the diffraction of radiation around the objects may be important even at visible wavelengths and naturally more important at longer wavelengths \\citep{Roques:6p1676}.  The discovery of more distant and more massive KBOs such as Eris \\citep{2005ApJ...635L..97B} begs the question of whether gravitational lensing of background stars by large KBOs and objects in the Oort cloud \\citep{Oort:1950p1427} is important. \\citet{2002astro.ph..9545C} argued that for distant massive KBOs and Oort cloud objects lensing may be important; furthermore, \\citet{2005ApJ...635..711G} argued that GAIA could measure the astrometric displacement from microlensing by an planet more massive than a few Jupiters within $10^4$~AU regardless of its location on the sky.  The technique in this paper probes much lower mass objects, but also exploits lensing to provide constraints on the properties of the asteroid.  Generally the diffractive effects of microlensing are neglected because the variation in the time delays across the lens is usually much larger than the coherence time of the observation, $1/\\Delta \\nu$ where $\\Delta \\nu$ is the bandwidth of the observation.  However, near a caustic crossing, diffraction may be important as argued by \\citet{1995ApJ...455..443J} to account for rapid variations in the light from Q2237+0305 (The Einstein Cross).  Typically the differential time delay highly magnified images for a point lens is about $2 GM/c^3$, the crossing time over the Schwarzschild radius of the lens; consequently, for the diffractive effects of lensing to be observable the Schwarzschild radius of the lens should be comparable or larger than the wavelength of the radiation.  The largest of the Kuiper Belt objects, Eris, has a Schwarzschild radius $R_S=2GM/c^2 \\approx 25 \\mu$m.  Therefore, quite naturally for observations of large KBOs diffractive microlensing may be important for observations in the near and mid-infrared whenever the gravitational effects are important.  This paper focuses on just this regime.  The first section, \\S\\ref{sec:diffr-micr}, outlines microlensing in the diffractive regime and generalizes the earlier results to include occultation.  This yields a expression for the transmission that is nearly identical to the unlensed result.  The next section, \\S\\ref{sec:results}, outlines the types of objects for which diffractive lensing may be important, examines several interesting cases and connects the diffraction patterns to the geometric limit (\\S\\ref{sec:geometric-optics}) in the limit where the lensing is weak. The final section (\\S\\ref{sec:conclusions}) outlines how diffractive lensing could constrain the properties of known objects and speculates on the probability of such lensing events. ",
        "conclusions": "\\label{sec:conclusions}  Microlensing is important for large asteroids (similar to Eris) in the Oort cloud and beyond; furthermore, in the near infrared and red-ward diffraction is important to understand the light curves from the combined microlensing and occultation of background stars by such objects.  The effects of microlensing are not covariant with variations in the velocity of source, lens and observer nor with variations in the impact parameter; therefore, observations of diffractive lensing combined with a measured distance to the asteroid constrain the mass and radius of the asteroid, or equivalently with an assumed value of the density of the asteroid, such observations would yield a mass, radius and distance.  Observations of diffractive microlensing may provide the only way of estimating the masses of such objects unless they have a satellite as Eris does \\citep{2007Sci...316.1585B}.  The angular size of even a large asteroid such as Eris is extremely small at the distance of the Oort cloud, so one would expect that the chance for an occultation would be tiny.  Over the course of a year, the asteroid will sweep out a region of the sky.  The chance of detection each year is proportional to this area.  The semimajor axis of the annual parallactic ellipse on the sky ranges from two to twenty arcseconds for a distance of $10^5$ and $10^4$~AU respectively.  The area of sky covered each year is given by the product of the arclength around the parallactic ellipse and angular diameter of the asteroid about 0.33~mas for an asteroid of radius 1200~km at $10^4$~AU; therefore, on average, at the nearer distance the asteroid would cover a region of sky 0.33 mas by 100 arcseconds or about 0.033 square arcseconds or a fraction $6 \\times 10^{-14}$ of the entire sky each year. The solid angle covered decreases inversely with the distance squared until microlensing starts to dominate, subsequently the area is proportional to $d^{-3/2}$.  It is safe to assume that the asteroid has sufficient velocity perpendicular to the Earth's orbital motion that it moves much more that one diameter per year (much greater than 8.2 cm/s); therefore, the asteroid will cover a new region each year. The asteroid itself may have a large proper motion that could increase the area further --- this motion is neglected in this analysis.  Any star within this area of sky will pass within the region $v<u_d$ in Figs.~\\ref{fig:1E5AU_Zero} through~\\ref{fig:1E4AU} over the course of that year.  If one assumes that the Oort cloud contains about ten Earth masses of objects similar to Eris at about $10^4$~AU, one would have to monitor about $4\\times 10^8$ stars for twelve years (twenty-four hours per day) to find a single microlensing occultation event by an object at $10^4$~AU.  The motion of the Earth in its orbit determines the duration of the event of about a minutes.  Currently the OGLE collaboration monitors about $4 \\times 10^8$ stars --- the present state of the art.  OGLE-III ran from 2001 to 2009 and made nearly $2 \\times 10^{11}$ photometric measurements \\citep{2008AcA....58...69U}.  The key is not only the sensitivity of the current slew of experiments, although monitoring additional stars or using several telescopes separated by at least the diameter of the asteroid would increase the detection rate, but the cadence of OGLE-III at about five minutes would simply undersample and miss such an event.  OGLE-IV will provide both a higher sensitivity and a higher cadence, but probably not a rapid enough cadence to distinguish such an event.  The best approach would be a difficult hybrid of the high cadence of observations such as \\citet{2006AJ....132..819R} on a large telescope and the monitoring of many stars as OGLE.  Because the goal is not small bodies as in \\citet{2006AJ....132..819R}, a 21ms-cadence is overkill; a one-second cadence would be sufficient allowing observations of much fainter stars.  Furthermore, stars with especially small angular sizes are not necessary.  Both of these factors along with the dedicated use of a large telescope could make probing the largest and presumably the rarest objects in the Oort cloud possible."
    },
    "0910/0910.0730_arXiv.txt": {
        "abstract": "{For Galactic open clusters, fundamental parameters such as age or reddening are usually determined independently of their integrated colours. For extragalactic clusters, on the other hand, they are derived by comparing their integrated colours with predictions of simple stellar population (SSP) models.} {We search for an explanation of the disagreement between the observed integrated colours of 650 local Galactic clusters and the theoretical colours of present-day SSP models.} {We check the hypothesis that the systematic offsets between observed and theoretical colours, which are $(B$$-$$V)\\approx 0.3$ and $(J$$-$$K_s)\\approx 0.8$, are caused by neglecting the discrete nature of the underlying mass function. Using Monte Carlo simulations, we construct artificial clusters of coeval stars taken from a mass distribution defined by an Salpeter initial mass function (IMF) and compare them with corresponding ``continuous-IMF'' SSP models.} {If the discreteness of the IMF is taken into account, the model fits the observations perfectly and is able to explain naturally a number of red ``outliers'' observed in the empirical colour-age relation. We find that the \\textit{systematic} offset between the continuous- and discrete-IMF colours reaches its maximum of about 0.5 in $(B$$-$$V)$ for a cluster mass $M_c=10^2\\,m_\\odot$ at ages $\\log t\\approx 7$, and diminishes substantially but not completely to about one hundredth of a magnitude at $\\log t >7.9$ at cluster masses $M_c> 10^5\\,m_\\odot$. At younger ages, it is still present even in massive clusters, and for $M_c \\leqslant 10^4\\,m_\\odot$ it is larger than 0.1 mag in $(B$$-$$V)$. Only for very massive clusters ($M_c>10^6\\,m_\\odot$) with ages $\\log t< 7.5$ is the offset small (of the order of 0.04 mag) and smaller than the typical observational error of colours of extragalactic clusters.} {} ",
        "introduction": "\\label{sec:intro}  Using data on accurate and homogeneous spatio-kinematic-photometric membership for 650 Galactic open clusters, we previously computed their integrated magnitudes in $B,V,J,H$, and $K_s$-passbands \\citep{intpar}. The magnitudes are based on accurate and uniform data from the \\ascc catalogue and were computed by adding the individual luminosities of the most secure cluster members. In contrast to previous lists of integrated magnitudes \\citep[e.g., those of][]{battin94, lata02} of Galactic star clusters, our data provide an independent and more importantly uniform dataset. Hence, our integrated magnitudes can and should be used as benchmarks for studies of the populations of extragalactic star clusters.  \\begin{figure}[t] \\resizebox{\\hsize}{36mm}{ \\includegraphics[clip]{13270f1a.ps} \\includegraphics[clip]{13270f1b.ps} }\\\\ \\resizebox{\\hsize}{36mm}{ \\includegraphics[clip]{13270f1c.ps} \\includegraphics[clip]{13270f1d.ps} } \\caption{Observed colours of 650 Galactic open clusters compared to theoretical colours computed from standard SSP models. The upper row is for $B$$-$$V$, the bottom one for $J$$-$$K_s$. The left column shows ``integrated colour $vs. \\log t\\,$'' diagrams, the right one the distributions of colour indices. Dots show Galactic open clusters from our sample. Black dots are young clusters used for constructing the colour distribution, grey dots the older clusters. The curves are the SSP model tracks (the solid one is from Starburst99, the dashed one is from GALEV).The observed colour distributions are shown as the filled cyan histogram, the GALEV model distribution as the hatched red one. } \\label{sev_fig} \\end{figure}  Based on these uniform data, we found disagreement between the observed colours of our cluster sample and theoretical colours derived from simple stellar population (SSP) models. The comparison of the \\citet{intpar} data with present-day SSP models such as GALEV \\citep{galev03} and Starburst99  \\citep{sb99_05} illustrated a substantial offset of the order of $0.2-0.3$ mag in $(B$$-$$V)$ of the model colours with respect to the observations for clusters younger than $\\log t \\lesssim 8.5$. This is both remarkable and important, since most data such as age and mass collected nowadays for the  bulk of extragalactic clusters, are derived by comparing their integrated light with the predictions of SSP models \\citep[see e.g.,][]{bikea03}. The models are produced by evolutionary synthesis codes simulating the extensive and complicated stellar populations of galaxies. However, they are also believed to be appropriate for describing cluster populations more than five orders of magnitude less massive (and apparently far simpler). Although star clusters may be influenced by effects that are irrelevant to galaxies (e.g., low number statistics, because of the discreteness of the real IMF, or dynamical mass-loss by evaporation of stars), the validity of the SSP approach for star clusters can be tested from general considerations \\citep{cervino04}. In contrast, in this paper we apply the issue of the discreteness of the IMF to the specific case of Galactic open clusters.  A colour offset of this kind was found previously by \\citet{lata02}, in the colour indices $U$$-$$B$, $B$$-$$V$, $V$$-$$R$, and $V$$-$$I$. However, they accept that only the offsets in $V$$-$$R$ and $V$$-$$I$ represent significant deviation of their observations from the model predictions. We suspect that the discreteness of the IMF in open clusters may explain these discrepancies. To study this issue in a more systematic way, we constructed our own SSP model that takes into account the effect of low number statistics (discreteness), which is a common problem for open star clusters. We cross-checked a continuous version of our code with GALEV and found good agreement.  The discreteness of the IMF itself is a natural assumption for any stellar ensemble consisting of individual objects. However, when one considers vast stellar populations (e.g., galaxies) that densely populate the entire colour-magnitude diagram, one can adopt a continuous IMF, which is more convenient for different technical reasons. For small populations (stellar clusters), the assumption of a discrete IMF is clearly appropriate.. This letter describes the solution of a restricted problem, namely the influence of the discrete-IMF approach on the colours of young clusters. The full scope of results related to the IMF-discreteness effect on star clusters will be reported elsewhere. ",
        "conclusions": "\\label{sec:conc}  Observations of local Galactic open clusters have enabled us to measure their integrated magnitudes, masses, ages, and reddening. Using this data set, we have found a remarkable discrepancy between the observed colours and the predictions of SSP models. The main reason for this disagreement is the neglect of the assumption of IMF-discreteness. When this effect is taken into account, the model agrees adequately with the observations and is even able to explain  the large colour spread observed in the empirical colour-age relation in a natural way.  In which conditions is the effect of discreteness relevant? Since it is a consequence of low-number statistics, it depends on the sparseness of the stellar population in the upper CMD of a cluster where the bulk of light is emitted. The density of the population depends primarily on cluster mass, and in addition on cluster age, on [strange at first glance] the lower mass limit of the stars formed, and to some degree on the slope of the IMF. In the context of this letter, other factors are not important.  According to our investigation, the \\textit{systematic} offset between the continuous- and discrete-IMF colours diminishes substantially but not completely at $\\log t >7.9$, at cluster masses $M_c> 10^5\\,m_\\odot$. At younger ages, it remains present even in massive clusters, and for $M_c \\leqslant 10^4$ it is larger than 0.1 mag in $(B$$-$$V)$. Only for very massive clusters ($M_c>10^6\\,m_\\odot$) and young ages ($\\log t< 7.5$), the offset falls below a typical observational error. We note in passing that the effect is stronger for redder passbands. These findings are in good agreement with the theoretical forecast of \\citet{cervino04}. The immediate consequence of the application of a continuous-IMF approach to the SSP models of stellar clusters is a systematic underestimate of reddening. Other problems affect the age and mass determination based on evolutionary grids of these models (at least for masses below $10^6\\,m_\\odot$). Because of the rather flat colour-age relation of young clusters, a systematic error of the order of 0.1 mag produces an error in cluster age of about one order of magnitude. On the other hand, there is a danger that the existence of a large number of ``main-sequence'' (i.e., blue) clusters could be interpreted as evidence of a recent burst of star formation. We can exclude such a burst having occured in our Galactic open clusters even if they are blue. Finally, the technique of the parameter determination should be changed to incorporate the cluster flux variations caused by the discrete nature of the upper IMF."
    },
    "0910/0910.5609_arXiv.txt": {
        "abstract": "}[2]{{\\footnotesize\\begin{center}ABSTRACT\\end{center} \\vspace{1mm}\\par#1\\par \\noindent {~}{\\it #2}}}  \\newcommand{\\TabCap}[2]{\\begin{center}\\parbox[t]{#1}{\\begin{center} \\small {\\spaceskip 2pt plus 1pt minus 1pt T a b l e} \\refstepcounter{table}\\thetable \\\\[2mm] \\footnotesize #2 \\end{center}}\\end{center}}   \\newcommand{\\TableSep}[2]{\\begin{table}[p]\\vspace{#1} \\TabCap{#2}\\end{table}}  \\newcommand{\\FigCap}[1]{\\footnotesize\\par\\noindent Fig.\\  % \\refstepcounter{figure}\\thefigure. #1\\par}  \\newcommand{\\TableFont}{\\footnotesize} \\newcommand{\\TableFontIt}{\\ttit} \\newcommand{\\SetTableFont}[1]{\\renewcommand{\\TableFont}{#1}}  \\newcommand{\\MakeTable}[4]{\\begin{table}[htb]\\TabCap{#2}{#3} \\begin{center} \\TableFont \\begin{tabular}{#1} #4 \\end{tabular}\\end{center}\\end{table}}  \\newcommand{\\MakeTableSep}[4]{\\begin{table}[p]\\TabCap{#2}{#3} \\begin{center} \\TableFont \\begin{tabular}{#1} #4 \\end{tabular}\\end{center}\\end{table}}  \\newenvironment{references}% { \\footnotesize \\frenchspacing \\renewcommand{\\thesection}{} \\renewcommand{\\in}{{\\rm in }} \\renewcommand{\\AA}{Astron.\\ Astrophys.} \\newcommand{\\AAS}{Astron.~Astrophys.~Suppl.~Ser.} \\newcommand{\\ApJ}{Astrophys.\\ J.} \\newcommand{\\ApJS}{Astrophys.\\ J.~Suppl.~Ser.} \\newcommand{\\ApJL}{Astrophys.\\ J.~Letters} \\newcommand{\\AJ}{Astron.\\ J.} \\newcommand{\\IBVS}{IBVS} \\newcommand{\\PASP}{P.A.S.P.} \\newcommand{\\Acta}{Acta Astron.} \\newcommand{\\MNRAS}{MNRAS} \\renewcommand{\\and}{{\\rm and }} {We report the analysis of a binary blue straggler in NGC 6752 with a short orbital period of 0.315 d and a W UMA-type light curve. We use photometric data spanning 13 years to place limits on the mass ratio ($0.15\\lesssim q \\lesssim 0.35$), luminosity ratio ($L_{1}/L_{2}\\approx 4.0$) and the ratio of the radii of the components ($r_{1}/r_{2}\\approx 2.0$). The effective temperatures of the components are nearly identical, and the system is detached or semi-detached (in the latter case the component filling its Roche lobe is the secondary). Such a configuration is unusual given the shortness of the orbital period, and it must have resulted from  substantial mass exchange. We suggest that some secondaries of W UMa-type stars, normally regarded as main sequence objects which fill their Roche lobes to different degrees, in fact may be shell-burning cores of originally more massive components. }  {\\bf Key words:} {\\it blue stragglers -- binaries: eclipsing -- globular clusters: individual (NGC~6752)} ",
        "introduction": "\\label{sect:intro}  Blue stragglers (BSs) occupy the extension of the main sequence above the turnoff point in color-magnitude diagrams of parent stellar populations. Since their detection in M3 (Sandage 1953) BSs have been identified in dozens of globular and open clusters. Two main mechanisms responsible for the formation of these objects have been advocated -- mass transfer in close binaries (McCrea 1964), and mergers resulting from direct collisions between stars in dense environments (Benz \\& Hills 1987). Knigge et al. (2009) found that the number of BSs belonging to a given globular cluster (GC) scales nearly linearly with the core mass of that cluster but is nearly independent of the core radius. This implies that the main channel leading to the formation of BSs is binarity. Moreover, Perets \\& Fabrycky (2009) argue that the progenitors of BSs form in primordial triple stars through the Kozai mechanism (Kozai 1962). This is in line with the observational evidence that most (possibly all) close binaries are formed in triple systems (Tokovinin et al. 2006).  The sample of candidate BSs in GCs includes several variables with light curves typical of contact binaries. The discovery of two W~UMa type systems in NGC~5466 (Mateo et al. 1990) was followed by the detection of further such systems in other clusters, mostly by the OGLE and CASE groups (see Rucinski 2000). Surprisingly, until now not a single mass determination or even a reliable light curve solution has been obtained for a GC contact binary. In fact, masses resulting from the analysis of light and radial velocity curves have been derived for just two binary BSs from GCs (Kaluzny et al. 2007a, 2007b). Some attempts have been made to estimate masses of apparently single BSs in GCs based on the modeling of spectra (\\eg De Marco et al. 2005) or pulsation periods of SX~Phe variables (Gilliland 1998). However, the actual accuracy of these determinations is hard to estimate.  The eclipsing binary V8-NGC~6752 (hereafter V8) was discovered by Thompson et al. (1999) during a photometric survey of the cluster field. The system has a W~UMa-type light curve, and an orbital period of 0.31~d. Our new photometry (Kaluzny \\& Thompson 2009) revealed that the secondary eclipse of V8 is total. As it was first discussed by Mochnacki \\& Doughty (1972), light curves of contact binaries showing total eclipses can be uniquely solved, yielding the mass ratio among other parameters. This is in contrast to contact systems with partial eclipses whose light curves usually cannot be reliably solved without spectroscopic information about the mass ratio.\\footnote{Except for relatively rare systems with mass ratios close to unity.} In short, the totality of the primary eclipse is sufficient to determine reliable apparent magnitudes of the individual components.  If V8 is indeed a contact binary belonging to the cluster then one can estimate its absolute parameters based solely on the photometric information, which together with the known distance to the cluster would yield the absolute magnitudes of its components. Then, based on effective temperatures estimated from colors, one can derive absolute radii of both stars. Knowing absolute radii and relative radii (the latter from the light curve solution) the absolute size of the orbit can then be derived,  yielding the total mass of the system from Kepler's law. Given the paucity of reliable determinations of masses for BSs in GCs we decided to analyze the light curves of V8 hoping to estimate its parameters. To our surprise it turned out that the binary is most likely a detached system with a mass ratio $q<0.35$. This is an unexpected result given the short orbital period of the system and the similar effective temperatures of its components. ",
        "conclusions": "\\label{sect:summary}  Our analysis indicates that, despite its short orbital period and W~UMa-type light curve, the binary V8 is a detached or semidetached system. This is an unexpected finding, as non-contact systems are extremely rare among non-degenerate binaries with periods $P\\leq 0.45$~d (Hilditch et al. 1988). To the best of our knowledge the only known systems of this kind are V361~Lyr (Hilditch et al. 1997; Kaluzny 1991; $P=0.30~d$) and OGLE-BW3-V38 (Maceroni \\& Rucinski 1997; $P=0.20$~d). The first of these is a semi-detached system with the more massive component filling its Roche lobe, while the latter is probably a detached system composed of two M dwarfs. Another possible non-contact binary of this kind is W~Crv (Rucinski \\& Lu 2000), however its actual configuration is poorly constrained. We note that V8 cannot be considered as an example of a system caught in the brief contact-break predicted by ``the thermal relaxation oscillations theory\" (Lucy \\& Wilson 1979). In such a system the {\\em primary} of V8 would have to fill its Roche lobe, and the components should have very different effective temperatures (the secondary being cooler). This is certainly excluded.   The stability and symmetry of the light curve of V8 together with the apparent constancy of the orbital period favors a detached configuration. Despite the totality of one of the eclipses the photometric analysis does not allow an accurate determination of the mass ratio of the system. However, it strongly favors the range $0.15\\lesssim q \\lesssim 0.35$. The components differ significantly in size ($r_{1}/r_{2}\\approx 2$), yet they have very similar effective temperatures. Figure 5 shows the location of the whole system and each of its components separately on the color-magnitude diagram of NGC~6752. When calculating the individual magnitudes of the components we adopted a luminosity ratio of 3.97, the value implied by the solution with $q=0.20$ for radiative envelopes. Note, however, that the derived value of $L_{1}/L_{2}$ depends rather weakly on the assumed $q$ (see Tables 2 and 3).  It is worth noting that the primary of V8 is located right at the extension of the unevolved main sequence of the cluster. This gives one more argument for the cluster membership of the binary, and suggests that the primary has the properties of an undisturbed main-sequence star. In contrast, the secondary is located far to the blue of the cluster main-sequence for its luminosity. Also, the observed luminosity ratio $L_{2}/L_{1}\\approx 1/4$ is much too high compared to that expected for main-sequence stars with $m_{2}/m_{1} < 0.35$. For stars with $0.43<m<2.0_{\\odot}$ we have $L\\sim m^{4}$ (Duric 2004). If the primary of V8 is indeed an ordinary main-sequence star then its mass can be estimated roughly at $0.8~m_{\\odot}$ given a cluster age of about 13~Gyr. In such a case one obtains $L_{2}/L_{1}=16$ for $q=0.5$, a conservative lower limit for V8, as in fact $q<0.35$. Clearly it is the the secondary which is the stranger component of the blue straggler V8.  Such an unusual system may throw some light on problems related to the evolution and observed status of W UMa-type binaries. The blind acceptance of the contact model (Lucy 1968a, 1968b) for very short period binaries with components of identical effective temperature, yet very different masses, has been recently strongly contested by the results for AW UMa (Pribulla \\& Rucinski 2008). This flagship case of the very low mass ratio binary ($q \\le 0.1$)  whose light curve so well obeys the contact model turned out not to be in contact upon a more careful spectroscopic scrutiny; instead, it appears to be a detached or semi-detached system engulfed in an optically thick stream with temperature not different from that of the stellar photosphere. The picture is indeed confusing, as the control case of V566 Oph ($q = 0.26$) turned out to agree -- both photometrically and spectrally -- with the contact model. Apparently, a range in conformance to the contact model exists.  There is no doubt that the present configuration of V8 results from a substantial mass exchange between the components (it is obvious that the secondary component is overluminous given the constraints on the present mass ratio). Most likely the secondary lost nearly all matter from its original envelope, and is burning hydrogen in a shell. Such a situation occurs in some low mass algols (Paczy\\'nski 1971; Eggleton \\& Kisleva-Eggleton 2002). However, before considering any detailed evolutionary scenario one needs to determine absolute parameters of the system. This requires supplementing available photometric data with spectroscopic observations. Obtaining a radial velocity curve for V8 is a challenging task. The luminosity ratio of $1/4$ is not a big obstacle, but the low apparent luminosity of the secondary ($V\\approx 18.8$) makes the observations  difficult as the short orbital period of the system limits the exposure time to 15-20 minutes. Useful spectra can be obtained only on the largest available telescopes.  The spectroscopic masses are essential in relating our current results to the important finding of Gazeas \\& Niarchos (2006) that the masses of secondary components in W UMa-type binaries are amazingly similar, clustering around 0.45 $m_\\odot$. It is possible that what we normally interpret as full-size secondaries (filling to different degrees their Roche lobes to different degrees) are in fact collapsed cores of originally more massive components,  as envisaged by Stepien (2006, 2009).   \\Acknow{ We sincerely thank Slavek Rucinski for his useful comments on the draft version of this paper. Research of JK, MR and KZ is supported by the grant MISTRZ from the Foundation for the Polish Science and by the grant N N203 379936 from the Ministry of Science and Higher Education. I.B.T. acknowledges the support of NSF grant AST-0507325. Some of the data presented in this paper were obtained from the Multimission Archive at the Space Telescope Science Institute (MAST). STScI is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS5-26555. Support for MAST for non-HST data is provided by the NASA Office of Space Science via grant NAG5-7584 and by other grants and contracts. }"
    },
    "0910/0910.4864_arXiv.txt": {
        "abstract": "{Cometary globule 12 is a relatively little investigated medium- and low mass star forming region 210 pc above the Galactic plane.} {This study sets out to discover the possibly embedded members of the CG~12 stellar cluster, to refine the NIR photometry of the known member stars and to study the star formation activity in CG~12 and its relation to the distribution of molecular gas, dust and mid- to far-infrared emission in the cloud.} { NIR \\J, \\H, and \\Ks \\ imaging and stellar photometry is used to analyse the stellar content and the structure of CG~12.} {Several new members and member candidates of the CG~12 stellar cluster were found. The new members include in particular a highly embedded source with a circumstellar disk or shell and a variable star with a circumstellar disk which forms a binary with a previously known A spectral type cluster member.  The central source of the known collimated molecular outflow in CG~12 and an associated ``hourglass''-shaped object due to reflected light from the source were also detected.  The maximum visual extinction in the cloud, based on observations of background stars, is $\\sim$20\\umag, but this is only a lower limit for the extinction through the two dense cloud cores. HIRES-enhanced IRAS images are used together with SOFI \\jhks \\ imaging to study the two associated IRAS point sources, 13546--3941 and 13547--3944. Two new 12\\,$\\mum$ \\ sources coinciding with NIR excess stars were detected in the direction of IRAS 13546--3941. The IRAS 13547--3944 emission at 12 and 25\\,$\\mum$\\ originates in  the Herbig~AeBe star h4636n and the 60 and 100\\,$\\mum$\\ emission from an adjacent cold source.} {} ",
        "introduction": "\\label{sec:introduction}  Low and intermediate mass star formation takes place primarily in isolated clusters (e.g., the CrA star forming cloud) and to a lesser extent in low mass star forming regions (e.g., Taurus and Chamaeleon).  As the stars form in dense gas/dust clouds, which are tightly concentrated towards the Galactic plane, star formation far out of the plane is not a likely event. Even though Chamaeleon and CrA are high Galactic latitude regions, they lie well within the scale height of molecular interstellar matter in the Galaxy (50-75 pc).  The well known nearby star forming regions have been studied in great detail from visual to radio wavelengths, but less is known of star formation regions off the Galactic plane.   \\citet{VanTilletal1975} detected strong CO emission in the direction of the reflection nebulosity \\object{NGC 5367} and, assuming a distance of 300 pc, calculated a mass of 30 \\Msun \\ for the mapped area.  NGC 5367 lies in the head of the Cometary Globule 12, \\object{CG 12}, listed in \\citet{HawardenBrand1976}.  Despite being classified by \\citet{HawardenBrand1976} as a cometary globule, together with those in the Gum nebula \\citep{reipurth1983}, CG 12 is actually a high-latitude low- and intermediate-mass star formation region.  With a Galactic latitude of 21\\degr \\ and at the distance of $\\sim$550\\,pc \\defcitealias{getmanetal2008}{GFLBG} (\\citet{maheswar1996} and \\citet{getmanetal2008}, hereafter GFLBG), CG~12 lies more than 200~pc above the plane.  A loose stellar cluster including at least three late B stars and two A stars \\citep{williams1977} is associated with CG~12. The associated reflection nebula  NGC 5367 is illuminated by the binary star \\object{h4636} (spectral types B4V and B7V). This binary contains at least one, or possibly two, Herbig~AeBe stars \\citep{ReipurthZinnecker1993}.  \\citetalias{getmanetal2008} find a high concentration of X-ray sources in the CG\\~12 region mainly in the direction of visible or NIR stellar objects and conclude that more than 50 of these are T-Tauri stars associated with the nebula.  \\citet{white1993} mapped the CG~12 centre region in CO (2--1) and \\ceo (2--1) and found a compact core and a highly collimated molecular outflow. Further molecular line observations of CG~12 are presented in \\citet{yonekuraetal1999a} and recently in\\defcitealias{haikalaolberg2007}{Paper 1} \\citet{haikalaolberg2007} \\citepalias{haikalaolberg2007} and\\defcitealias{haikalaetal2006}{Paper 2} \\citet{haikalaetal2006} \\citepalias{haikalaetal2006}. The structure and dynamics of the CG~12 core have been discussed in detail in Papers 1 and 2.  The molecular material lies in a 10$'$ NS elongated lane with three compact \\ceo maxima, \\cgnp, \\cgsp, and \\cgswp. The southern maximum, \\cgsp, lies near the binary h4636.  The \\ceo emission in \\cgs traces mainly warm gas on the surface of a dense core (called \\dcoplp \\ core) revealed in \\dcopl and CS emission and in 1.2\\,mm continuum \\citepalias{haikalaolberg2007, haikalaetal2006}. The centre of the highly collimated molecular outflow reported in \\citet{white1993} lies in the direction of the \\dcopl core.  The northern core, \\cgnp, is cold and the relative velocities and intensities of \\ceo and high density tracers indicate that molecular material is highly depleted. The molecular mass of the CG~12 core area, as traced by the \\ceo observations, is larger than 100 \\Msun. This mass corresponds only to the area observed in \\ceo and traced by the \\ceo emission. The mass of the cloud envelope not visible in \\ceo and that outside the mapped area is much larger, as shown by e.g. the large scale CO mapping of \\citet{yonekuraetal1999a} according to which the mass of the globule is larger than 300 \\Msun.  The linear size of the CG~12 core region ($>$1 pc) and the molecular mass are of the same order as those of known nearby low mass star formation regions like Chamaeleon~I.  CG~12 has been recently reviewed by \\citet{2008hsf2.book..847R}.  The median photometric age of the T-Tauri population in CG~12 is $\\sim$4 Myr, but the apparent spread of ages is considerable, from 1 to 20 Myr \\citepalias{getmanetal2008}.  The more than 50 T-Tauri stars detected by \\citetalias {getmanetal2008} represent possibly the first generation of stars born in the CG~12 area, and the B and A stars \\citep{williams1977} may be the second generation.  The four IRAS point sources and the collimated molecular outflow in the CG~12 cloud show that star formation is presently going on.  Because of the great distance and the high extinction towards the cloud core, the limiting magnitudes of the Two Micron All Sky Survey (2MASS) survey allow only the detection  of the brightest stars embedded in the cloud in all  three \\jhks \\ colours.  A deeper \\J,\\H,\\Ks \\ survey of CG~12 than 2MASS is clearly needed to study the possible embedded stellar population. Spitzer NIR IRAC imaging data on CG~12 has recently become public and can also be used to analyse the embedded stars.  We report imaging of the core of CG~12 in \\J, \\H,  and \\Ks \\ NIR bands with SOFI at the NTT telescope at La Silla. Observations, data reduction, and calibration procedures are described in Sect.~\\ref{sec:observations} and the observational results in Sect~\\ref{sec:results}. The new results are discussed and compared with available data at other wavelengths in Sect.~\\ref{sec:discussion}.  The results are summarised and the conclusions drawn in Sect.~\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  The head of the cometary globule CG~12 has been observed in the \\J, \\H, and \\Ks \\ bands.  The new NIR photometry combined with already existing data at other wavelengths allows the analysis of the star formation and the content of the associated stellar cluster in more detail than was possible before.  Star formation in CG 12 (within the region imaged in \\jhks \\ with SOFI) is concentrated on and near the two strong \\ceo maxima \\cgs and \\cgnp. Apart from the known members of the associated stellar cluster several new, embedded cluster members were detected as well. Because the SOFI images are much deeper than the available 2MASS survey data it is possible to define more accurately the NIR colours of the already known embedded member stars which were near the 2MASS limit.  Apart from the stellar photometry, the SOFI imaging allows the study of the cloud surface brightness due to scattered light. Scattered light permits us to pinpoint the central source of the highly collimated molecular outflow detected by \\citet{white1993} and to detect the probable shadow cast by the circumstellar disk around star \\four.  Using NIR \\J,\\H,\\Ks \\ photometry alone is not in all cases sufficient to decide whether some stars are indeed members of the CG~12 stellar cluster. In some cases an associated nebulosity can confirm the membership. For four stars the observed X-ray activity \\citepalias{getmanetal2008} indicates that the stars are likely members.   The conclusions are the following:  1. Seven embedded (proto)stars (\\one, \\two, \\three, \\four, \\five, \\six\\ and W77-6b) are identified.  The visual extinction of these stars ranges from 10 to more than 30 magnitudes.  The stars \\three \\ and \\four \\ were already known members. Stars \\two, \\ \\three, \\ \\four \\ and \\five \\ are also  detected in X-rays \\citepalias{getmanetal2008}.  2. Stars \\one, \\four, \\five, \\six, and W77-6b have infrared excess indicating a circumstellar shell or disk. The shadow of the probable disk around star \\four \\ is seen in the cloud surface brightness.  3. Scattered light from the probable central source of the collimated outflow detected by \\citet{white1993} as well as an hourglass-shaped cavity in the parental molecular cloud formed by the outflow are detected. Apart from  the hourglass  a larger scale cavity cleared by the outflow to the less tenuous part of the molecular cloud is also seen. The stellar source is better resolved in the Spitzer IRAC images at 3.4, 4.5 and 5.8 $\\mum$ \\ than in the SOFI \\H \\ and \\Ks \\ images, where it blends into the bright background nebulosity.   {\\bf 4.} The \\citet{williams1977} star W77-6 is a binary with a separation of $\\sim$4\\arcsec.  The brighter component, W77-6a is an early A type star with a visual extinction of $\\sim$3 magnitudes. The fainter component,  W77-6b, has  NIR excess indicating  circumstellar material. The W77-6b \\J \\ magnitude changed by 0\\fm6 between two observing nights. The Spitzer NIR imaging suggests that a third source and an associated nebulosity not visible in the \\jhks \\  imaging may be present.  5. HIRES enhanced IRAS 12 and 25\\,$\\mum$ \\ maps combined with IRAS fluxes and the NIR imaging suggest that the IRAS point source 13547--3944 consists of two separate sources. The 12 and 25 $\\mum$ \\ emission originates in  the h4636n component and the longer wavelengths from an adjacent, embedded cold cloud source.  The 12 $\\mum$ \\ HIRES map of IRAS 13546--3941 resolves into two previously undetected sources coinciding with the binary stars W77-6a/b and \\three/\\four. The HIRES 25 $\\mum$ maximum lies between the 12 $\\mum$ \\ sources near the IRAS 13546--3941 nominal position and the \\ceo maximum and a mm continuum source. The 25\\,$\\mum$\\ detection could be a source embedded in the continuum core.  6. The SOFI data are not deep enough to probe the visual extinction through the two dense cloud cores in the direction of \\cgn and \\cgsp. The observed maximum extinction around \\cgn is $\\sim$20\\umag, and $\\sim$15\\umag \\ around \\cgsp. No background stars are seen through the core centres."
    },
    "0910/0910.3733_arXiv.txt": {
        "abstract": "% Image compression has been a frequent topic of presentations at ADASS. Compression is often viewed as just a technique to fit more data into a smaller space.  Rather, the packing of data -- its ``density'' -- affects every facet of local data handling, long distance data transport, and the end-to-end throughput of workflows.  In short, compression is one aspect of proper data structuring.  For example, with FITS tile compression the efficient representation of data is combined with an expressive logistical paradigm for its manipulation. \\par A deeper question remains.  Not just how best to represent the data, but which data to represent.  CCDs are linear devices.  What does this mean?  One thing it does not mean is that the analog-to-digital conversion of pixels must be stored using linear data numbers (DN). An alternative strategy of using non-linear representations is presented, with one motivation being to magnify the efficiency of numerical compression algorithms such as Rice. ",
        "introduction": "The Rice compression algorithm (Rice {\\em et\\ al.} 1993) is particularly familiar to astronomers in combination with the FITS Tiled Image Convention (White {\\em et\\ al.} 2006, Seaman {\\em et\\ al.} 2007). It has been informally paired in the past with non-linear data representations (Nieto-Santisteban {\\em et\\ al.} 1999, Nicula {\\em et\\ al.} 2005). Mention of compression is often followed immediately by the word ``scheme'', as if this branch of computer science were suitable only for black-box heuristics arrived at through a process of trial-and-error.  This paper seeks to begin a process of layering the issue of optimal data encoding onto a more formal foundation.  As discussed in Pence, Seaman \\& White (2009A,B), the compression achieved for an astronomical image is typically determined almost entirely by its background noise.  Any processing that quiets the background will increase the compression ratio.  This is the foundation for all lossy compression algorithms -- in effect, lossy compression combines a preprocessing step that tempers an image's noise characteristics with the subsequent application of a lossless algorithm.  Janesick (2001) describes a non-linear hardware or software component called a ``square-rooter''.  Instead of using a linear analog-to-digital conversion when reading out a CCD, the square-root of these values acts to linearize the Poisson statistics that characterizes these detectors.  Such a DN encoding has been used on several space missions to reduce bandwidth requirements for data transport. While technically lossy if the square-root transform occurs after the A/D conversion, it is equally reasonable to consider this a type of lossless encoding since the transform maintains oversampling of the noise at both the high and low end of the dynamic range.  Square-root encoding itself acts as a form of compression -- 65,636 linear data levels (for 16-bit pixels) will turn into a scant 256 levels after the square-root.  This is not nearly as dramatic when expressed in bits, since mapping 16 bits into 8 bits corresponds to a compression ratio of just $R = 2$.  The more important aspect, rather, is to linearize the noise. ",
        "conclusions": ""
    },
    "0910/0910.4113_arXiv.txt": {
        "abstract": "We study the spectroscopic properties and environments of red (or passive) spiral galaxies found by the Galaxy Zoo project. By carefully selecting face-on, disk dominated spirals we construct a sample of truly passive disks (\\ie ~they are not dust reddened spirals, nor are they dominated by old stellar populations in a bulge). As such, our red spirals represent an interesting set of possible transition objects between normal blue spiral galaxies and red early types, making up $\\sim6\\%$ of late-type spirals. We use optical images and spectra from SDSS to investigate the physical processes which could have turned these objects red without disturbing their morphology. We find red spirals preferentially in intermediate density regimes. However there are no obvious correlations between red spiral properties and environment suggesting that {\\it environment alone is not sufficient to determine if a galaxy will become a red spiral}.  Red spirals are a very small fraction of all spirals at low masses ($M_\\star < 10^{10}$\\msun), but are a significant fraction of the spiral population at large stellar masses showing that {\\it massive galaxies are red independent of morphology}.  We confirm that as expected, red spirals have older stellar populations and less recent star formation than the main spiral population.  While the presence of spiral arms suggests that major star formation cannot have ceased long ago (not more than a few Gyrs), we show that these are also not recent post-starburst objects (having had no significant star formation in the last Gyr), so {\\it star formation must have ceased gradually}. Intriguingly, red spirals are roughly four times as likely than the normal spiral population to host optically identified Seyfert/LINER (at a given stellar mass and even accounting for low luminosity lines hidden by star formation), with most of the difference coming from objects with LINER-like emission. We also find a curiously large optical bar fraction in the red spirals ($70\\pm5\\%$ verses $27\\pm5\\%$ in blue spirals) suggesting that {\\it the cessation of star formation and bar instabilities in spirals are strongly correlated}. We conclude by discussing the possible origins of these red spirals. We suggest they may represent the very oldest spiral galaxies which have already used up their reserves of gas - probably aided by strangulation or starvation, and perhaps also by the effect of bar instabilities moving material around in the disk. We provide an online table listing our full sample of red spirals along with the normal/blue spirals used for comparison. ",
        "introduction": "The advent of large galaxy surveys like the Sloan Digital Sky Survey (SDSS) in which photometry (and therefore colours) are readily available for millions of objects has lead to the common use of optical colours to define ``early\" and``late\" type galaxy samples \\citep[e.g.][]{S09,S08,LP07,C07,B06,C05}. This method is particularly favoured since obtaining morphologies for large numbers of galaxies has until recently been impossible. This simplification is justified since it has been shown many times that the majority of galaxies follow a strict colour-morphology relation. For example \\citet{M09} argued that 85\\% of galaxies to z$\\sim$1 are either red, bulge-dominated galaxies or blue, disk dominated galaxies; while \\citet{C06} showed a similar result for 22000 low redshift galaxies (both using automated methods for morphological classification).  However the clear correlation between colour and morphology is surprising, given that the colours of galaxies are determined primarily by their stellar content (and therefore their recent star formation history, mostly within the last Gyr) while the morphology is primarily driven by the dynamical history. The clear link between colour and morphology then gives a strong indication that the timescales and processes which drive morphological transformation and the cessation of star formation are strongly related - at least in most cases. In this paper however, we consider a class of object (the red spirals) where the link described above appears to be broken.  Since the morphology-density relation was first quantified \\citep{D80} many mechanisms have been proposed for the transformation of blue, star forming, disk galaxies in low density regions of the universe, to red, passive, early type galaxies in clusters. A recent review of many of the proposed mechanisms, and the evidence supporting them, can be found in \\citet{BG06}. Clearly two things must happen for a star forming blue spiral galaxy to turn into a passive red early type. First, star formation must cease (which can indirectly alter the morphology by causing spiral arms and the disk in general to fade, possibly producing an S0 or lenticular from a spiral), and secondly, in order to produce a {\\it bona fide} elliptical the same, or a different process must also dynamically alter the stellar kinematics of the galaxy.  The presence of an unusually red or passive (\\ie. non-star forming) population of spiral galaxies in clusters of galaxies was first noticed by \\citet{V76} in the Virgo cluster. Later studies of distant cluster galaxies in {\\it Hubble Space Telescope} (HST) imaging also revealed a significant number of so-called ``passive'' spiral galaxies with a lack of on-going star formation (Couch et al. 1998; Dressler et al. 1999; Poggianti et al. 1999).  Passive late-type galaxies were identified at lower redshifts in the outskirts of SDSS clusters by \\citet{G03}, using concentration as a proxy for morphology. Passive spirals in a cluster at $z\\sim0.4$ were studied by \\citet{M06} who found star formation histories from GALEX observations consistent with the shutting down of star formation from strangulation (as described by \\citealt{Bekki02}). Passive spirals have also been revealed in a cluster at $z\\sim 0.1$ in the STAGES survey, using HST morphologies \\cite{W09}, rest frame NUV-optical SEDs \\citep{W05} and 24$\\mu$m data from Spitzer \\citep{G09}. In that series of papers, ``dusty red late-types\" and ``optically passive late-types\" are found to be largely the same thing, with a non-zero (but significantly lowered) star formation rate revealed by the IR data.  Red spirals/late types have been studied in several recent papers \\citep{Lee08,D09,HC09,CH09}, as well as \\citet{MR09} who talk about blue passive galaxies (\\ie ~ galaxies with blue colours, but showing no indication of recent star formation in their spectra) which mostly appear to have late type morphologies and have very recently shut down star formation. These might be the progenitors of the red spirals. \\citet{Bundy09} have studied the redshift evolution of red sequence galaxies with disk like components in COSMOS and use it to estimate that as many as 60\\% of spiral galaxies must pass through this phase on the way to the red sequence - making it an important evolutionary step.  The Galaxy Zoo project \\citep{L08} revealed the presence of a significant number of visually classified spiral galaxies which are redder than the blue cloud (between 16-28\\% of the total galaxy population depending on environment, \\citealt{B09}). In this paper we study in more detail the physical properties and environments of this population of red spiral galaxies. Galaxies drawn from this population have the morphological appearance of spiral galaxies with a distinct spiral arm structure, but have rest-frame colours which are as red as a typical elliptical galaxy, indicating little or no recent star formation activity. We are studying these objects in order to identify the physical process which is most important in their formation.  It is clear that all spiral galaxies can be affected by various physical processes as they evolve -- in this paper we attempt to identify which are most important for red spirals, asking how they are able to shut down star formation while retaining their spiral morphology. A list of possible mechanisms includes processes that depend on environment such as: (1) galaxy-galaxy interactions: in high density regions there is an increased probability of interaction with other galaxies. Most major mergers destroy spiral structure \\citep{TT72} unless they involve very gas rich progenitors \\citep{H09}, but some interactions can be quite gentle (\\eg~ Walker, Mihos \\& Hernquist 1996), for example minor-mergers, tidal interactions etc. (2) Interaction with the cluster itself also occurs and can remove the gas which forms the reservoir for star formation. This can be due to tidal effects \\citep[e.g.][]{Gn03}, or interaction with the hot intercluster gas, either through thermal evaporation \\citep{CS77} or ram pressure stripping \\citep{GG72}. (3) Processes like harrassment \\citep{M99} and starvation or strangulation \\citep{L80,Bekki02} have also been shown to have a significant effect on late type galaxies. Harrassment refers to the heating of gas by many small interactions, while starvation or strangulation refers to the gradual exhaustion of disk gas after the hot halo has been stripped away. These mechanisms both occur at much larger cluster radii (\\ie ~lower densities) than the \"classic\" environmental effects. Internal mechanisms could be more important.  For example, (4) the latest semi-analytical models of galaxy formation all invoke feedback from a central massive black hole (or active galactic nuclei; AGN) to explain the most massive red elliptical galaxies \\citep{G04,S05,Sc06,Cr06,Bower06}, although the effect of this process on disk galaxies has been studied less, it still may have some effect (Okamoto, Nemmen \\& Bower 2008). (5) Another culprit could be bar instabilities in spiral galaxies which drive gas inwards (eg. Combes \\& Sanders 1981) and may trigger AGN activity and/or central star formation (eg. Shlosman et al. 2000), perhaps using up the reservoir of gas in the outer disk and making spirals red. (6) Finally red spirals could simply be old spirals which have used up all their gas in normal star formation activities without having any major interactions. In normal spirals, the gas that feeds ongoing star formation comes from infall of matter from a reservoir in the outer halo \\citep{BG06}. As first suggested by Larson, Tinsley \\& Caldwell (1980) and expanded by \\citet{Bekki02} the removal of gas from this outer halo (``strangulation\", or ``starvation\") will cause a gradual cessation of star formation proceeding over several Gyrs.  We describe the sample and data along with the selection of the ``red spirals\"  in more detail in Section 2. In Section 3 we discuss the stellar populations and star formation history of the red spirals. In Section 4 we discuss their environment and the environmental dependence of star formation. The impact of AGN is considered in Section 5, and bar instabilities are discussed in Section 6. In Section 7 we discuss the plausible mechanisms for formation of the red spirals and future directions which could be taken to distinguish between them. We present a summary and conclusions in Section 8. The adopted cosmological parameters throughout this paper are $\\Omega_{m}$ = 0.3, $\\Omega_{\\Lambda}$ = 0.7 and $H_{0} = 70$ \\kmsMpc. ",
        "conclusions": "We study the interesting population of red, or passive, spiral galaxies found by the Galaxy Zoo project. We identify from these red spirals a population of intrinsically red true disk dominated spirals by limiting the sample in inclination (to reduce the impact of dust reddening), requiring that spiral arms be visible, and removing bulge dominated systems using the SDSS parameter {\\tt fracdeV}. We compare this sample to blue spirals selected in the same way and find that:  \\begin{itemize} \\item{Red spirals are more likely to be more massive, luminous galaxies than blue spirals. They represent an insignificant fraction of the spiral population at masses below $10^{10}$\\msun ~but are significant at the highest masses, showing that {\\it massive galaxies are red regardless of morphology}.}  \\item{At the same mass as blue spirals, face-on red spirals do not have larger amounts of dust reddening (as measured by the Balmer decrement), therefore their red colours indicate an ageing stellar population not an increased dust content.}  \\item{Red spirals have lower (but not zero) rates of ongoing and recent star formation when compared to blue spirals. This is partly related to their higher average mass, however {\\it at a fixed mass, red spirals still have less recent star formation than blue spirals}.}  \\item{As previously observed, red spirals are more common at intermediate local densities (around, or just inside the infall regions of clusters). They are also observed to be more likely than blue spirals to have close neighbours.}  \\item{Red spirals in all environments have lower rates of recent and on-going star formation than blue spirals, and there are no significant trends of the star formation rates with environment when spirals are split into red/blue. Clearly the process which creates red spirals is not confined to regions of high galaxy density. So {\\it environment alone is not sufficient to determine whether a galaxy will become a red spiral or not}.}  \\item{Red spirals are more than four times more likely to be classified as Seyfert+LINER/composite objects from their optical spectra than blue spirals. This is partly due to the higher masses of red spirals but is still observed when they are compared to a blue spiral sample selected to have the same mass distribution. We find that a small fraction of low luminosity AGN are being revealed as the star formation is turned off in the red spirals, but this is not enough to account for all the difference. Most of the difference comes from an increased fraction of LINER-like emission ($82\\pm12\\%$ of Seyfert/LINERs found in red spirals are LINERs compared to $57\\pm7\\%$ in blue spirals).}  \\item{Red spirals have significantly higher bar fractions than blue spirals (70\\% versus 27\\%), suggesting that bar instabilities and the shutting down of star formation in spirals are correlated}. \\end{itemize}  We propose three possible origins for the red spiral population studied in this work and suggest the most likely explanation is that a combination of the three accounts for the shutting down of their star formation while they retain their spiral structure: \\begin{enumerate} \\item{Perhaps red spirals are just old spirals which have used up all of their gas. They are found preferentially in intermediate density regions because structures first starts to form at the peaks of the dark matter distribution, but in the centres of clusters spiral morphologies cannot stand up to the environmental disturbances. Red spirals then represent the end stages of spiral evolution irrespective of environment (and in the absence of major mergers) - the spiral version of ``downsizing\".} \\item{Perhaps red spirals are satellite galaxies in massive dark matter halos. In this scenario, they are accreted onto the halo as a normal blue spiral and have experienced either strangulation (where the gas in their outer halos has been gently stripped off, and no further cold gas has been allowed to accrete) or harassment (heating their disk gas and preventing further star formation). Low mass spirals would probably be disrupted in this process and so are not observed as red spirals.} \\item{Perhaps red spirals evolved from normal blue spirals which had bars that were particularly efficient at driving gas inwards. This removed gas from the outer disk and turned the spiral red. If it triggered star formation in the central regions it must have occured more than $\\sim 1$ Gyr ago since red spirals are not post starburst galaxies.}  \\end{enumerate}  The red spirals in this work probably cannot be the progenitors of S0s as they have a significantly higher bar fraction than in observed in the S0 population. S0s may however be the end product of red spirals with larger bulges than we have studied here.  \\paragraph*{ACKNOWLEDGEMENTS}  This publication has been made possible by the participation of more than 160,000 volunteers in the Galaxy Zoo project. Their contributions are individually acknowledged at \\texttt{http://www.galaxyzoo.org/Volunteers.aspx}.  KLM acknowledges funding from the Peter and Patricia Gruber Foundation as the 2008 Peter and Patricia Gruber Foundation International Astronomical Union Fellow, and from the University of Portsmouth and SEPnet (www.sepnet.ac.uk). Support for the work of MM in Leiden was provided by an Initial Training Network ELIXIR (EarLy unIverse eXploration with nIRspec), grant agreement PITN-GA-2008-214227 (from the European Commission). AKR, MM, HCC, RCN acknowledge financial support from STFC. Support for the work of KS was provided by NASA through Einstein Postdoctoral Fellowship grant number PF9-00069 issued by the Chandra X-ray Observatory Center, which is operated by the Smithsonian Astrophysical Observatory for and on behalf of NASA under contract NAS8-03060. CJL acknowledges support from The Leverhulme Trust and the STFC Science In Society Programme. Funding for the SDSS and SDSS-II has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Science Foundation, the U.S. Department of Energy, the National Aeronautics and Space Administration, the Japanese Monbukagakusho, the Max Planck Society, and the Higher Education Funding Council for England. The SDSS Web Site is http://www.sdss.org/."
    },
    "0910/0910.3213_arXiv.txt": {
        "abstract": "Using 11-years of OGLE V-band photometry of Q2237+0305, we measure the transverse velocity of the lens galaxy and the mean mass of its stars. We can do so because, for the first time, we fully include the random motions of the stars in the lens galaxy in the analysis of the light curves. In doing so, we are also able to correctly account for the Earth's parallax motion and the rotation of the lens galaxy, further reducing systematic errors. We measure a lower limit on the transverse speed of the lens galaxy, $v_{\\rm t} > 338\\kms$ (68\\% confidence) and find a preferred direction to the East. The mean stellar mass estimate including a well-defined velocity prior is $0.12 \\leq \\left<M/M_\\sun\\right> \\leq 1.94$ at $68\\%$ confidence, with a median of $0.52~M_\\sun$. We also show for the first time that analyzing subsets of a microlensing light curve, in this case the first and second halves of the OGLE V-band light curve, give mutually consistent physical results. ",
        "introduction": "\\label{sec:intro}  Quasar microlensing provides a unique tool for studying the properties of cosmologically distant lens galaxies and the structure of quasars \\citep[see][]{Wambsganss06}. Each of the multiple images of the quasar passes through the gravitational potential of the stars along the line-of-sight in the lens galaxy. These stars microlens each of the ``macro'' images, so the total magnification of each quasar image is strongly affected by the lensing effects of the stars and the size of the quasar emission region. Since the observer, lens galaxy, stars, and source quasar are all moving, these magnifications change on timescales of 1-10 years with order unity amplitudes.  The relevant physical scale for quasar microlensing is the Einstein radius projected into the source plane plane, \\begin{eqnarray} \\rE &=& D_{\\rm OS} \\sqrt{\\frac{4G\\left<M\\right>}{c^2}\\frac{D_{\\rm LS}}{D_{\\rm OL}D_{\\rm OS}}} \\nonumber \\\\ &=& 1.8\\times 10^{17} \\left(\\frac{\\left<M\\right>}{M_\\sun}\\right)^{1/2}\\cm, \\end{eqnarray} where $G$ is the gravitational constant, $c$ is the speed of light, $\\left<M\\right>$ is the mean stellar mass of the stars, $D_{\\rm LS}, D_{\\rm OL}$, and $D_{\\rm OS}$ are the angular diameter distances between the lens-source, observer-lens, and observer-source respectively, where we have used the lens and source redshifts for Q2237+0305 \\citep[$z_{\\rm l}=0.0394\\,,z_{\\rm s}=1.685$,][Q2237 hereafter]{Huchra85}. The observed microlensing amplitude is controlled by the ratio between the source size, $R_{\\rm V} \\approx 6\\times10^{15}\\cm$ (in V-band, see our companion paper \\citet{Poindexter09}, hereafter Paper II) and $\\rE$, in the sense that smaller accretion disks produce larger variability amplitudes than larger disks. If a source is much larger than $\\rE$, there is little change in the magnification.  The timescale for variability is determined by the relative velocities of the observer, the lens, its stars, and the source.  Generally, the lens motions dominate \\citep{Kayser89}, leading to two characteristic timescales. There is the timescale to cross an Einstein radius, \\begin{eqnarray} \\tE &\\approx& \\frac{\\rE}{v_{\\rm lens}} \\frac{(1+z_{\\rm l})\\dol}{\\dos} \\nonumber \\\\ &\\approx& 8 \\left(\\frac{v_{\\rm lens}}{462\\kms}\\right)^{-1} \\left(\\frac{\\rE}{2\\times10^{17}\\cm}\\right) {\\rm yr}, \\nonumber \\\\ \\label{eqn:te} \\end{eqnarray} and there is the timescale to cross the source, \\begin{eqnarray} \\ts &\\approx& \\frac{R_{\\rm V}}{v_{\\rm lens}} \\frac{(1+z_{\\rm l})\\dol}{\\dos} \\nonumber \\\\ &\\approx& 0.4 \\left(\\frac{v_{\\rm lens}}{462\\kms}\\right)^{-1} \\left(\\frac{R_{\\rm V}}{6\\times10^{15}\\cm}\\right) {\\rm yr}, \\nonumber \\\\ \\label{eqn:ts} \\end{eqnarray} where $v_{\\rm lens}$ is the expected transverse speed of the lens for Q2237. These timescales are also affected by the direction of motion relative to the shear \\citep{Wambsganss90}. Variation is guaranteed on timescale $\\tE$ and can be observed on timescale $\\ts$ if the magnification pattern locally has structure on the scale of $R_{\\rm V}$.  Quantitative studies using quasar microlensing have exploded in the last few years. Recent efforts have studied the relationships between accretion disk size and black hole mass \\citep{Morgan10}, size and wavelength \\citep{Anguita08,Bate08,Eigenbrod08b,Poindexter08,Floyd09,Mosquera09}, sizes of non-thermal (X-ray) and thermal emission regions \\citep{Pooley07,Morgan08a,Chartas09,Dai09}, and the dark matter fraction of the lens \\citep{Dai09,Pooley09,Mediavilla09}. All these studies have used static magnification patterns which ignore the random stellar motions in the lens galaxy. However, the stellar velocity dispersions of lens galaxies are comparable to the peculiar velocities of galaxies, and ignoring them can lead to biased results. For example, the average direction of motion of the source is the same for all images, but this coherence is limited by the random motions of the stars.  With fixed patterns one must either overestimate the coherence, by using the same direction of motion for each image, or underestimate it by using independent directions for each image \\citep[see][]{Kochanek07}. In either case, it would be dangerous to attempt measurements that depend on this coherence, such as disk shape and orientation, or the transverse peculiar velocity of the lens.  \\citet{Kundic93}, \\citet{Schramm93}, and \\citet{Wambsganss95} considered the effects of random stellar motions and found that these motions can also lead to shorter microlensing time scales because the pattern velocities of the microlensing caustics can be much higher than any physical velocities. As a result, measurements based on static magnification patterns may underestimate source sizes or mean masses or overestimate the transverse velocity in order to match the effects created by the random stellar motions. \\citet{Wyithe00} showed that it is possible to statistically correct for these effects and that the velocity correction can be up to 40\\% depending on the optical depth and shear. Another benefit to dynamic magnification patterns is the ability to properly account for the velocity of Earth around the Sun and the rotation of the lens galaxy \\citep{Tuntsov04}.  Dynamic magnification patterns are also important because they impart a well-measured physical scale to the patterns. All direct microlensing observations are in ``Einstein units'' where the length scale is $\\left<M\\right>^{1/2}\\cm$. Determining masses, velocities, or source sizes requires some sort of dimensional prior.  In our studies we have generally followed \\citet{Kochanek04a} and used velocity priors designed to mimic the combined effects of random and ordered motion.  The reason for focusing on the velocity is that we know two of the contributions, our velocity and the stellar velocities, and the remaining peculiar velocities of the lens and source are truly random variables for which we have reasonable priors from cosmological models. Source sizes turn out to be little affected by the choice of priors \\citep[see the discussion in][]{Kochanek04a}, but estimates of the true velocity and the mean stellar masses are affected.  Hopefully by including the true random stellar motions we can further reduce the sensitivity of microlensing results to such priors.  \\begin{figure*} \\plotone{tracks} \\caption{Example of a trial source trajectory (dark line segments) superposed on an instantaneous point-source magnification pattern for $\\mmass = 0.3$. Darker shades indicate higher magnification. An HST H-band image in the center labels the images and the corresponding magnification patterns. Each pattern is rotated to have the correct orientation relative to the lens. This particular LC2 trial has an effective lens-plane velocity of $\\sim 600\\kms$ Northeast. The solid disk at right has a radius of $10^{17}\\cm$ on these patterns.} \\label{fig:tracks} \\end{figure*}  In this paper we use microlensing to measure the peculiar velocity of a lens galaxy and the mean mass of its stars including the effects of the stellar motions, Earth's motion, and the rotation of the lens galaxy. The transverse velocity direction can be measured with microlensing because the shear sets a preferred direction for each image and the statistics of variability depend on the motion relative to this axis (see Figure \\ref{fig:tracks}). In theory, accurately measuring the transverse peculiar velocity of many galaxies over a broad range of redshifts could form the basis of a new cosmological test \\citep{Gould95}. Measuring the mean stellar mass, including remnants, is an independent means of checking local accountings \\citep[e.g.][]{Gould00}, which must be assembled from very disparate selection methods for high mass, low mass, evolved and dead, remnant stars. Moreover, doing this is possible in detail only for the Galaxy. While microlensing is relatively insensitive to the mass function \\citep[see][]{pac86,Wyithe00b}, there are some prospects of exploring this in the future as well \\citep[e.g.][]{Wyithe01,Schechter04,Congdon07}.  This work expands on the methods described in \\citet{Kochanek04a} and \\citet{Kochanek07} by adding the random stellar motions in the lens galaxy. In this paper we address the computational issues and then apply this improved technique to determine the transverse motion of Q2337 and the mean mass of its stars. In Paper II, we study the shape of the accretion disk of the source quasar. We describe the photometric data in \\S\\ref{sec:data}. Then we describe the Bayesian Monte Carlo Method and the models we use in \\S\\ref{sec:methods}. Our results are presented in \\S\\ref{sec:results} followed by a discussion in \\S\\ref{sec:discussion}. We use an $\\Omega_M = 0.3, \\Omega_{\\Lambda} = 0.7$ flat cosmological model with $H_0 = 72~\\rm km~s^{-1}~Mpc^{-1}$. ",
        "conclusions": "\\label{sec:discussion}  By including the random motions of the stars, we can now use microlensing to study the peculiar velocity of the lens galaxy and to estimate mean stellar masses and potentially the stellar mass functions with fewer systematic uncertainties. In particular, we find a clear preference for the direction of motion of the lens galaxy. In fact, as we use a less restrictive velocity prior, the direction of motion is better constrained since faster speeds are allowed. We cannot however, determine the speed without additional priors. If we assume a mean stellar mass of $\\left<M\\right> = 0.3 M_\\sun$, we find that the peculiar velocity is $v_{\\rm t} < 486 \\kms$ which is consistent with the other estimates \\citep{Wyithe99,GilMerino05} but more fully includes all the physical uncertainties. It should not be surprising that we can determine a dimensionless quantity, the direction of motion, better than the dimensional speed, given the basic problem of microlensing that all observables are in $\\sqrt{\\left<M\\right>}\\cm$. Very roughly (see Figure \\ref{fig:tracks}), the preferred direction has images A and B moving more closely to perpendicular to the ridges of the magnification patterns created by the shear and images C and D are moving more parallel to the shear direction. This is consistent with the variability of A/B compared to C/D. We included the parallax effects of the Earth's orbit, and the results weakly favored its inclusion.  We had hoped that modeling the random motions would be more of a help in breaking these degeneracies by setting a physical scale. This is probably true for low mean masses $\\left<M\\right>$. For fixed variability amplitudes, reproducing the light curves with a low mean mass requires small physical velocities, while high mean masses require high velocities. Adding the stellar motions at their observed dispersion eliminates low mass solutions independent of the unknown peculiar velocities by setting a floor to the velocity scale. High mass solutions need peculiar velocities, $\\sigma$, that are larger than the stellar motions, $\\sigma_*$, and so are only constrained by the priors on the peculiar velocities. Essentially, the dynamic patterns act like static patterns once $\\sigma_* < \\sigma$, and we recover the familiar degeneracies of static patterns. Thus, our correct treatment of the stellar motions constrains low mass but not high mass solutions in the absence of a peculiar velocity prior. With a well-defined cosmological prior on $\\sigma$ \\citep{Tinker09}, we find that $0.12 \\leq \\left<M/M_\\sun\\right> \\leq 1.94$ at $68\\%$ confidence, demonstrating that the microlensing objects are typical of stellar populations and their remnants. This mass range is consistent with expectations for normal stellar populations (see \\S\\ref{sec:mass}), but not tightly constraining.  We largely ignore the macro magnifications predicted by the mass distribution of the lens galaxy in our calculations because of their systematic uncertainties. However some recent studies have made use of this information by analyzing image pairs straddling a critical curve which should have the same magnification \\citep{Floyd09,Bate08}. A concern is that the macro magnification may be affected by undetected substructure, differential extinction, or contamination by the lens or host galaxy. In our standard analysis we use the AC signal and largely discard the DC signal by not tightly constraining the mean magnification. Given sufficiently long light curves, the results will converge to the true magnification offsets.  Even for Q2237, with its decade long OGLE light curves, the data are not sampling long enough paths across the patterns (see Figure \\ref{fig:tracks}) to show convergence.  At present, the distribution of differential mean magnification offsets are too broad (Figure \\ref{fig:zero}) to tightly constrain any systematic magnification offsets. Fortunately, our results for the other physical parameters are little affected by whether we allow these offsets to vary or constrain them with the extinction estimates of \\citet{Agol09}.  Finally, we show for the first time that microlensing variability in a lens gives the same results when analyzing different portions of its light curve. The analysis of light curves LC1 and LC2, corresponding to the 1st and 2nd halves of the 11 year OGLE monitoring period, lead to statistically consistent distributions for every parameter we consider. This both confirms our ability to measure parameters and gives us tighter constraints after combining the results. It would be computationally challenging to analyze the full light curve simultaneously because it becomes (exponentially?) harder to fit longer light curves. However, such full analyses are likely needed for some quantities, particularly the magnification offsets, to converge.  In the future we will likely include binaries, even though their effect is not likely to influence the results other than interpreting the meaning of the mean stellar mass (by up to $0.05M_\\sun$, as discussed in \\S\\ref{sec:results}). However, like the projection of our motion relative to the CMB, the streaming velocities in Q2237 are small compared to the peculiar velocities, and so are do little to break the degeneracy. The effects of streaming velocities will be seen most strongly in true disk lenses (none are known, except, potentially PMN~J2004$-$1349, \\citet{Winn01}), or in lenses such as Q~J0158$-$4325 (see \\citet{Morgan08b} for a microlensing analysis of this active system) lying close to the equator of the CMB dipole, which will have the full $369\\kms$ dipole motion \\citep{Hinshaw09}. These CMB equatorial lenses should also show significantly shorter microlensing variability time scales. Detecting this effect would be an independent confirmation of the kinematic origin of the dipole.  Q2237 was a natural first candidate for a full analysis with moving stars because of the excellent OGLE data, short microlensing timescales, and negligible time delays between the images. However, there is no problem extending our approach to analyzing microlensing data with moving stars to any other microlensing analysis.  Even if the time delays are unknown, cases with different trial delays could simply be tried sequentially \\citep{Morgan08b}. Moreover our method can easily be extended to multi-wavelength data sets to examine the structure of the accretion disk varies with wavelength \\citep{Poindexter08}. The memory requirements would be too great to fit each band simultaneously as in \\citet{Poindexter08}, but we can use a modified version of the method \\citet{Dai09} applied to the joint optical and X-ray analysis of RXJ1131$-$1231. The models are first run on the band with the most epochs. As good fitting trials are found, the starting points, velocities, and $\\chi^2$ matrix are saved.  Next, for each successive band, we recompute the light curves corresponding to the epochs and source sizes of the other wavelengths, and the results of these new fits are used to continue the $\\chi^2$ calculation. Since the overall execution times are only modestly longer than using static stars, there is no reason not to use this more physically correct approach."
    },
    "0910/0910.3449_arXiv.txt": {
        "abstract": "{ The spectrum of cosmic rays (CRs) is affected by their escape from an acceleration site. This may have been observed not only in the gamma-ray spectrum of young supernova remnants (SNRs) such as RX~J1713.7$-$3946, but also in the spectrum of CRs showering on the Earth. } {The escape-limited model of cosmic-ray acceleration is studied in general. We discuss the spectrum of CRs running away from the acceleration site. The model may also constrain the spectral index at the acceleration site and the ansatz with respect to the unknown injection process into the particle acceleration. } { Analytical derivations. We apply our model to CR acceleration in SNRs and in active galactic nuclei (AGN), which are plausible candidates of Galactic and extragalactic CRs, respectively. In particular, for young SNRs, we take account of the shock evolution with cooling of escaping CRs in the Sedov phase. } {The spectrum of escaping CRs generally depends on the physical quantities at the acceleration site, such as the spectral index, the evolution of the maximum energy  of CRs and the evolution of {{\\bf the normalization factor of the spectrum. }} It is found that the spectrum of run-away particles can be both softer and harder than that of the acceleration site. } { The model could explain spectral indices of both Galactic and extragalactic CRs produced by SNRs and AGNs, respectively, suggesting the unified picture of CR acceleration. } ",
        "introduction": "\\label{sec:intro} The origin of cosmic rays (CRs) has been one of the long-standing problems. The number spectrum of nuclear CRs observed at the Earth, $\\mathcal{N} (E) \\propto E^{-s}$, shows a break at the ``knee'' energy ($\\approx10^{15.5}$~eV), below which the spectral index is about $s \\approx -2.7$ \\citep{cronin1999}. Because of the energy-dependent propagation of CRs, the spectral shape at the source is different from that observed at the Earth. Taking into account the propagation effect, the source spectral index has been well constrained as $s \\approx 2.2-2.4$ in various models \\citep[e.g.,][]{strong98,putze2009}. This value of $s$ has been also inferred in order to explain the Galactic diffuse gamma-ray emission \\citep[e.g.,][]{strong00}. This fact may give us valuable insights on the acceleration mechanism of CRs.  Mechanisms of CR acceleration have also been studied for a long time and the most plausible process is a diffusive shock acceleration (DSA) \\citep{krymsky77,axford77,bell78,blandford78}. Very high-energy gamma-ray observations have revealed that the existence of high-energy particles at the shock of young supernova remnants (SNRs), which supports the DSA mechanism as well as the paradigm that the Galactic CRs are produced by young SNRs \\citep[e.g.,][]{enomoto02, aharonian04, aharonian05c, katagiri05}. Recent progress of the theory of DSA has revealed that the back-reactions of accelerated CRs are important if a large number of nuclear particles are accelerated \\citep{drury81,malkov01}. There are several observational facts which are consistent with the predictions of such nonlinear model \\citep{vink03, bamba03b, bamba05a, bamba05b, warren2005, uchiyama07, helder2009}. The model predicts, however, the harder spectrum of accelerated particles at the shock than $f (p) \\propto p^{-4}$ (corresponding to $s=2$) where $p$ is the momentum of CRs, in particular, near the knee energy \\citep{malkov97, berezhko99, kang01}. This fact apparently contradicts with the source spectral index of $s \\approx {2.2-2.4}$ inferred from the CR spectrum at the Earth. Even in the test-particle limit of DSA, such a soft source spectrum requires a shock with the small Mach number ($M \\la10$), which is unexpected for young SNRs.  There are several models of DSA, depending on the boundary conditions imposed. Different models predict different spectra of CRs dispersed from the shock region. So far, the age-limited acceleration has been frequently considered as a representative case (\\S~\\ref{sec:agelimit}). In this model, all the particles are stored around the shock while accelerated. When the confinement becomes inefficient, all the particles run away from the region at a time. Then, the source spectrum of CRs which has just escaped from the acceleration region is expected to be the same as that at the shock front. Therefore, this model predicts that the source spectrum is the same as that of accelerated particles, which is typically harder than the observed one. In this paper, we consider an alternative model, the escape-limited acceleration, to explain the observed CR spectrum at the Earth (\\S~\\ref{sec:escape}). This model is preferable to the age-limited acceleration when we consider observational results for young SNR RX~J1713.7$-$3946, of which TeV $\\gamma$-ray emission is more precisely measured than any others (\\S~\\ref{sec:rxj1713}).   The nature of CRs with energies much higher than the knee energy is also still uncertain. While CRs below the second knee ($\\sim {10}^{17.5}$~eV) may be Galactic origin, the highest energy CRs above $\\sim{10}^{18.5}$~eV are believed to be extragalactic. Possible candidates are active galactic nuclei (AGNs) \\citep[e.g.,][]{BS87,Tak90,RB93,PMM09}, gamma-ray bursts \\citep[][]{Wax95,Vie95,MINN06}, magnetars \\citep[][]{Aro03,MMZ09} and clusters of galaxies \\citep[][]{KRB97,ISMA07}. The intermediate energy range from $\\sim {10}^{17.5}$~eV to $\\sim{10}^{18.5}$~eV is more uncertain. Both the Galactic and extragalactic origins are possible and it may just a transition between the two. As the extragalactic origin, AGNs \\citep{berezinsky06}, clusters of galaxies \\citep{MIN08} and hypernovae \\citep{WRMD07} have been proposed so far.  Among these possibilities, AGN is one of the most plausible candidates for accelerators of high-energy CRs, because it can explain the UHECR spectrum above $\\sim 10^{17.5}$~eV assuming the proton composition. In such a proton dip model, the source spectrum of UHECRs is $\\mathcal{N} \\propto E^{-s}$ with $s=2.4$--2.7, depending on models of the source evolution \\citep{berezinsky06}. The required source spectral index of $s=2.4$--2.7 can be explained by several possibilities. First, it can be attributed to the acceleration mechanism itself. One can consider non-Fermi acceleration mechanisms \\citep{berezinsky06} or the two-step diffusive shock acceleration in two different shocks \\citep{Alo+07}. Second, the index can be attributed to a superposition of many AGNs with different maximum energies, and one can suppose that AGNs with different luminosities may have different maximum energies \\citep{KS06}. Recently, \\citet{berezhko2008} proposed another possibility under the cocoon shock model. In this cocoon shock scenario, different maximum energies can be interpreted as maximum energies of escaping particles at different ages of AGN jets. Although it is very uncertain whether efficient CR acceleration occurs there, this scenario would also be one of the possibilities to be investigated in detail.   The organization of the paper is as follows. After the brief introduction of the age-limited model of the CR acceleration (\\S~\\ref{sec:agelimit}), we study the escape-limited model in \\S~\\ref{sec:escape}. For a simple understanding, the general argument in a stationary, test-particle approximation is given in \\S~\\ref{app:testparticle}. Then, we derive the formulae of the maximum energy of accelerated particles in \\S~\\ref{sec:pmax_escape}, and of the spectrum of escaping particles in \\S~\\ref{sec:spect_escape}. We consider the applications to young SNR and AGN in \\S~\\ref{sec:snr} and \\S~\\ref{sec:AGN}, respectively. Section~\\ref{sec:discussion} is devoted to a discussion. ",
        "conclusions": "\\label{sec:discussion}  In this paper, we have investigated the escape-limited model of CR acceleration, in which the maximum energy of CRs of an accelerator is limited by the escape from the acceleration site. The typical energy of escaping CRs decreases as the shock decelerates because the diffusion length becomes longer. After revisiting the escape-limited model and reconsider its detail more generally, we have derived a simple relation between the spectrum of escaping particles and one in the accelerator. Then, using the obtained relation, we have discussed which model of injection is potentially suitable to make the Galactic and extragalactic CRs observed at the Earth.  For young SNRs, we have considered the shock evolution with cooling by escaping CRs and those spectra for the three injection models. As a result, we have found that in the case PH, it is difficult to satisfy the condition for the source spectrum of Galactic CRs ($s_{\\rm esc} \\approx 2.2$--2.4). On the other hand, $s_{\\rm esc} \\approx 2.2$--2.4 can be achieved in cases of PS and TL.  We have also applied our escape-limited model to AGN cocoon shocks as well as young SNRs. This model is just one of the various candidates proposed so far, even if AGNs are UHECR accelerators. Nevertheless, it is interesting that the young SNR and the AGN cocoon shock scenarios can explain the Galactic and extragalactic cosmic rays observed at the Earth in the same picture for all the three injection models if we accept the proton-dip model inferring $s_{\\rm esc} \\approx2.4$--2.7. Whether the proton dip model is real or not can also be tested by future UHECR and high-energy neutrino observations. In this paper, we have focused on the proton case. Obviously, heavier nuclei become important above the knee so that we need to take into account them in order to explain the CR spectrum over the whole energy range. We can also apply the escape-limited model to heavy nuclei CRs for this purpose, although it is beyond the scope of this paper.  We point out a potential problem for the magnetic field amplification in the escape-limited model. In the case of young SNRs, we have determined the evolution of the maximum energy  in the phenomenological way, and adopted $\\alpha \\approx 6.5$. Using Eqs.~(\\ref{eq:pmax_esc2}) and (\\ref{eq:pm_general}), we obtain $B \\propto u_{\\rm sh}^{\\frac{10+2c}{3}}$, where $\\ell \\propto R_{\\rm sh}^c$. The same result is obtained for  AGN cocoon shocks because we have considered the case in which the same parameters describe both the young SNR shocks and the AGN cocoon shocks. In particular, for $c =1$ as is used in Eq. (31), we obtain $B^2 \\propto u_{\\rm sh}^8$ which means that $B$ rapidly decreases with radius (or time). In principle, both the dependence of $B$ on $u_{\\rm sh}$ and the value of $c$ can  be determined theoretically, and then the evolution of  the maximum energy should be predicted. Some previous works are based on theoretical arguments on the magnetic field evolution \\citep[e.g.,][]{bell04,PLM06,berezhko2008,caprioli09}, which seem to be different from our phenomenological one. At present, the mechanisms of the particle acceleration and the magnetic field amplification are still highly uncertain despite of many theoretical efforts \\citep[e.g.,][]{niemiec08,requelme09, ohira09a, LM09}. Hence, we expect that further theoretical and observational studies might reveal this discrepancy or exclude the possibility of escape-limited acceleration in the future.   In this paper, we have mainly considered spectra of dispersed CRs around young SNRs and AGN cocoon shocks. However, applications to other astrophysical objects are, of course, possible. For example, the old SNRs detected by {\\it Fermi} LAT, such as W28, W44, W51 and IC443 \\citep{abdo09,abdo09w51c} have been of great interest because they likely generate escaping CRs. In fact, the number of CRs around such old SNRs is likely to decrease with time or the shock radius, that is $\\beta<0$ while $\\alpha>0$. For example, when we consider the dynamics of an old SNR, we have $u_{\\rm sh} \\propto t^{-2/3} \\propto R_{\\rm sh}^{-2}$ \\citep[e.g.,][]{yamazaki06}, so that we have $E \\propto u_{\\rm s}^2R_{\\rm s}^3\\propto R_{\\rm sh}^{-1}$, i.e., $\\delta = -1 < 0$. On the other hand, the value of $\\alpha$ may be different from 6.5, which could be attributed to various complications such as the interaction with the dense molecular cloud, and so on. For example,  for the maximum hardening case, that is, $s_\\mathrm{esc}=s + (s-1)(\\sigma^{-1}-1)/\\alpha$ (see Appendix A.2), we find $s_{\\rm esc} \\approx 1.5$ when $\\alpha \\approx 1$ and $s \\approx 2.5$ where we assume $s=(\\sigma +2)/(\\sigma -1)$. This might be the case for old SNRs such as W51C \\citep{abdo09w51c}. In addition, the maximum energy may be rather small for the old SNRs, so that the spectrum above $p_m^{(\\rm esc)}$ would be suppressed. The spectrum of high-energy gamma rays might give us important information on both the acceleration and escape processes of CRs with energies much lower than the knee energy."
    },
    "0910/0910.2353_arXiv.txt": {
        "abstract": "We present the results of simulations of shadows cast by a zone plate telescope which may have one to four pairs of zone plates. From the shadows we reconstruct the images under various circumstances. We discuss physical basis of the resolution of the telescope and demonstrate this by our simulations. We allow the source to be at a finite distance (diverging beam) as well as at an infinite distance (parallel beam) and show that the resolution is worsened when the source is nearby. By reconstructing the zone plates in a way that both the zone plates subtend the same solid angles at the source, we obtain back high resolution even for sources at a finite distance. We present simulated results for the observation of the galactic center and show that the sources of varying intensities may be reconstructed with accuracy. Results of these simulations would be of immense use in interpreting the X-ray images from recently launched CORONAS-PHOTON satellite. ",
        "introduction": "\\label{intro}  Zone Plate Telescopes (ZPTs) have generated immense theoretical interest in the past. (Ables, 1968; Dicke, 1968; Barrett \\& Swindell, 1996; Desai, Norris \\& Nemiroff, 1993; Desai et al. 1993; 1998, 2000). Due to the availability of technology to make finer zones using X-ray and gamma-ray opaque materials, the interest of using zone plates have increased recently. In Chakrabarti et al. (2009, hereafter Paper I), we presented extensive theoretical studies and some of the results of the experiments conducted at the Indian Centre for Space Physics X-ray laboratory on zone plate telescopes. Such telescopes have been used in Indian payload system RT-2 aboard the Russian satellite CORONAS-PHOTON which was launched on 30th January, 2009. As mentioned in Paper I the ZPTs have an advantage over other high resolution X-ray telescopes in that they can have arbitrarily high angular resolution and that the resolution can also be independent of energy bands in a large range of energy. The only disadvantage is of course that the ZPT  is a two element system as opposed to the conventional coded aperture masks (CAMs) which are single element systems.  In the present paper of our series we shall present the results of the Monte-Carlo simulations of the nature of photon distribution on the detector planes for various combinations of zone plates and X-ray sources. The basic starting point in this paper would be Paper I and references therein. Here, we consider one pair to four pairs of zone plate combinations, with a point source and distributed sources. We also vary the distance of source from finite to infinite. We hope to be able to present the most comprehensive results in this subject so that interpretation of the data from CORONAS-PHOTON may become easier. In our next paper (Nandi et al. 2009) the instrumentation and early results from RT-2 payload which uses ZPTs would be presented.  The plan of the present paper is the following: In the next Section, we discuss the angular resolution of a ZPT and also the variation of the nature of the reconstructed image when the source distance varies. In Section 3, we will present in detail the results of our simulations in various circumstances. Finally, in Section 4, we make concluding remarks. ",
        "conclusions": "In this Paper, we presented results of various numerical simulations obtained by varying the source, the zone plate telescope parameters and the detector parameters. We considered both two-pair and four-pair configurations which enabled us to remove the pseudo-source and both the pseudo-source and DC-offset respectively. We pointed out that the pseudo-source cannot be totally removed if the source itself is at a finite distance. This is due to the fact that at a finite distance, different pair is hit by photons from the source at various angles and the contributions to the pseudo-source do not cancel out. We showed that the resolution of the instrument is worsened when the source is placed at a finite distance. However, we showed that if the telescope is modified in a way that both the plates subtend equal angles at the source, the high resolution is recovered. Of course, a practical difficulty of such a method is that the second zone plate has to be modified dynamically as the source distance is varied, though, for instance for medical purposes, one could keep the source at a fixed distance always and thus a fixed sized ZP2 would suffice. Alternatively, one could convolve the distorted fringes obtained due to sources at a finite distance by theoretical Moir\\'{e} pattern obtained for the source at infinite distance before deconvolving the pixel information. This is beyond the scope of the present paper and will be dealt with elsewhere.  We studied the cases with both CMOS and CZT detectors and showed that the setup with CZT detectors will have very limited field of view in order to even reconstruct a single source. This is because of the large pixel size. With a CMOS detector we simulated how a zone plate telescope would view the center of the Galaxy. When there are multiple sources of varying intensity, we showed that one could first obtain a general image and then improve upon it by subtracting the fringes from the strongest ones.  The zone plate telescopes have very high potential for future space astronomy, especially since it will have higher resolutions over a large range of energy. They may also be used for medical science since the sources at finite distances can also be resolved well if modified ZPTs are used. Recently ZPTs have been used in RT-2/CZT payload aboard Russian satellite CORONAS-PHOTON for the first time. The satellite has been recently launched and the results are expected in near future, especially when the sun becomes active.  The results of the zone plate imagers will be discussed elsewhere (Nandi et al. 2009)."
    },
    "0910/0910.5059_arXiv.txt": {
        "abstract": "{} {Recently a new analysis of cluster observations in the Milky Way found evidence that clustered star formation may work under tight constraints with respect to cluster size and density, implying the presence of just two sequences of young massive cluster. These two types of clusters each expand at different rates with cluster age.} {Here we  investigate whether similar sequences exist in other nearby galaxies.} {We find that while for the extragalactic young stellar clusters the overall trend in the cluster-density scaling is quite comparable to the relation obtained for Galactic clusters, there are also possible difference. For the LMC and SMC clusters the densities are below the Galactic data points and/or the core radii are smaller than those of data points with comparable density. For M83 and the Antenna clusters the core radii are possibly comparable to the Galactic clusters but it is not clear whether they exhibit similar expansion speeds. These findings should serve as an incentive to perform more systematic observations and analysis to answer the question of a possible similarity between young galactic and extragalactic star clusters sequences.} {} ",
        "introduction": "\\begin{figure}[t] \\resizebox{\\hsize}{!}{\\includegraphics[angle=-90]{Pfalzner_fig1.ps}} \\caption{{Cluster density as a function of cluster size for clusters more massive than 10$^3$ \\Msun as published by Pfalzner (2009). references therein.}} \\label{fig:cluster_rad} \\end{figure}   Lada \\& Lada (2003) showed that in our Galaxy most stars do not form in isolation but in cluster environments. These young clusters typically consist of thousand stars or more with densities ranging from $<$0.01 to several 10$^5$ \\Msun pc$^{-3}$ The wide variety of observed cluster densities lead to the assumption that clusters are formed over this entire density range. Albeit Maiz-Apellaniz (2001) noticed the existance of two types of clusters in the Galaxy, only recently Pfalzner (2009) found that massive clusters develop in a bimodal way as in Fig.\\ref{fig:cluster_rad}, where the red, green and blue symbols represent clusters with ages $t_c <$ 4 Myr,  4 Myr $ < t_c <$ 10 Myr,  10 Myr $< t_c <$ 20 Myr, respectively. Two well-defined sequences in the density-radius plane emerge showing the {\\em bi-modal nature of the cluster evolution.}  \\begin{figure}[t] \\resizebox{\\hsize}{!}{\\includegraphics[angle=0]{LMC_size_density.eps}} \\caption{a) Cluster density as a function of cluster radius and b) cluster size as a function of cluster age. Here the LMC and SMC (purple circle) clusters are shown as filled symbols and the Milky Way as open symbols (same clusters as in Fig.~\\ref{fig:cluster_rad}). The values for the clusters in the LMC are taken from Mackey \\& Gilmore 2003a and McLaughlin \\& van der Marel 2005, the SMC value from Mackey \\& Gilmore 2003b.} \\label{fig:LMC_cluster_size_age} \\end{figure}  Pfalzner (2009) classified the two modes as starburst and leaky cluster sequences. The starburst cluster sequence implies a population of compact clusters (0.1 pc) with high initial densities (10$^5$ - 10$^6$ \\Msun pc$^{-3}$) which then expand with the mass-density decreasing as $\\sim R^{-3}$ (an actual $2\\sigma$-fit gives a $R^{\\alpha}$-dependence with $\\alpha$=-2.71 $ \\pm $ 0.32) and evolve to have sizes of a few pc over a period of 10 Myr or longer. Prominent members of this type of cluster are for example Arches, NGC 3603 and Westerlund 1. The leaky cluster sequence implies the creation of a second population of diffuse clusters ($\\sim$ 5pc) which expand loosing mass during the process, until they have sizes of a few tens of pc. The expansion also follows a predefined track, but here the density decreases as $\\sim R^{-4}$ (actual fit value $\\alpha=-$4.07$\\pm0.33$). Here NGC 6611, Ori 1a-c or U Sco are typical examples. Note, the cluster radii were not determined exactly the same way for all clusters shown in Fig.~\\ref{fig:cluster_rad}. However, star burst cluster values all represent the core radii apart from $\\chi$ Per and h Per (for a discussion on the determination of these two radii see  Pfalzner 2009).  It follows that star formation occurs only under an extremely limited set  of conditions, and may require a fundamental revision of star formation hypotheses in our Galaxy.  This immediately raises the questions of the origin of these two distinct cluster sequences and whether similar density-radius correlations could be found in other galaxies. This Letter addresses the latter question.   Most extragalactic clusters one observes are likely to belong to the starburst cluster sequence because although their extension is on average smaller the luminosity of their high number of O-stars is much easier to detect than the smaller number of O stars spreading over a larger area in leaky clusters. So in all figures subsequent to Fig.~1 only the starburst clusters of the Galaxy are shown for comparison and the radial extent is the core radius.   \\begin{table*} \\begin{center} \\begin{tabular}{l|*{5}{l}} \\hline \\\\[-2ex] Galaxy & radius determination   &  mass determination &  background  &  age determination\\\\\\hline \\\\[-2ex] LMC$^1$          &  King profile fit        & from mass/light ratio from    & mean surface brightness     & combination of IMF $^8$ \\\\ &  half-surface brightness & evolutionary code$^5$         & in annulus at r $\\gg R_c$   & and evolutionary code$^5$  \\\\ SMC$^2$          &  King profile fit        & from mass/light ratio from    & mean surface brightness     & Based on color-magnitude  \\\\ &  half-surface brightness & evolutionary code$^5$        & in annulus at r $\\gg R_c$    & diagram$^{12}$  \\\\ Antennae$^3$     &  Fit with ISHAPE code$^6$    & UBVI-colors and mass/light   &   Fit with ISHAPE code      & LICK line strength indices  \\\\ &                          & ratio$^7$ with extinction-   &                             & and $^{11}$  \\\\ &                          & corrected SSP   &                             &                        \\\\ M 83$^4$         &  Fit with ISHAPE code$^6$    & Isochromes$^9$  &   Fit with ISHAPE code      & UVBI-colors, 2-color and  \\\\ &                          &                               &                          & S-sequnce diagram $^{10}$ \\end{tabular} \\caption{Determination methods of radius, mass and age and background consideration ($^1$Maxkey \\& Gilmore 2003a, $^2$ Mackey \\& Gilmore 2003b, $^3$ Mengel et al. 2008, $^4$ Larsen \\& Richter 2004, $^5$Fioc \\& Rocca-Volmerangen 1997, $^6$ Larsen 1999, $^7$Anders \\& Fritz-von-Alvensleben 2003, $^8$ Kroupa et al. 1993,  $^9$ Girardi 2000, $^{10}$ Girardi 1995, $^11$ Schweizer 2004 $^{12}$ Chiosi 1995. \\label{table:method}} \\end{center} \\end{table*}   Currently the Milky Way is not in a intense star formation phase accounting for the scarcity of starburst clusters in the Galaxy. Another reason is that extinction at low latitudes hampers the detection of distant Galactic clusters. By contrast, there are starburst and interacting galaxies such as the \"Antennae\", where many clusters with masses $>$ 10$^4$ \\Msun\\ are observed with ages $<$1Gyr (Zhang \\& Fall 1999). In between, there exist the dwarf starburst galaxies such as NGC 4214, where several massive clusters with ages $<$100 Myr are visible in the central regions.  This Letter shows a first investigation of whether similar development tracks can be found for clusters in such different galaxies. Observationally one is restricted to nearby galaxies, and even there radii are only resolved for a very limited sample. Here we scanned the literature for clusters younger than 30 Myr in the Large Magellanic Cloud (LMC), Small Magellanic Cloud (SMC), Messier 83 (M83) and the Antennae. ",
        "conclusions": "For the extragalactic young stellar clusters the overall trend in the cluster radius versus cluster density relation seems to fit the relation obtained for Galactic starburst clusters by Pfalzner (2009). A possible difference is indicated for the LMC clusters, where the densities are mostly below the comparable Galactic data at equal radius, or the radii are smaller than those of data points with comparable density. This might indicates that the sizes for LMC clusters are systematically smaller than the Galactic cluster sizes. This could be explained by intrinsic observational effects.    However, the sample size of extragalactic young clusters ($<$ 20 Myr ) with known radii and masses is so small that this investigation can only be regarded as a first hint that the young starburst cluster sequence might also exist in galaxies other than the Milky way. In this study only literature values of age, mass and radius were used mostly obtained by different methods. Further studies with a larger sample size and more homogenous data acquisition will be required to confirm this result."
    },
    "0910/0910.5868_arXiv.txt": {
        "abstract": "\\noindent Timing noise in the data on accretion-powered millisecond pulsars (AMP) appears as irregular pulse phase jumps on timescales from hours to weeks. A  large systematic phase drift is also observed in the first discovered AMP  \\sax.  To study the origin of these timing features, we use here the data of the well studied 2002 outburst of \\sax.  We develop first a model for pulse profile formation accounting for the screening of the antipodal emitting spot by the accretion disk.  We demonstrate that the variations of the visibility of the antipodal spot associated with the receding accretion disk cause a systematic shift in Fourier phases, observed together with the changes in the pulse form. We show that a strong secondary maximum can be observed only in a narrow intervals of inner disk radii, which explains the very short appearance of the double-peaked profiles in  \\sax.  By directly fitting  the pulse profile shapes with our model, we find that the main parameters of the emitting spot such as its mean latitude and longitude as well as the emissivity pattern change irregularly causing small shifts in pulse phase, and the strong profile variations are caused by the increasing inner disk radius. We finally notice that significant variations in the pulse profiles in the 2002 and 2008 outbursts of \\sax\\ happen at fluxes differing by a factor of 2, which can be explained if  the inner disk radius is not a simple function of the accretion rate, but depends on the previous history. ",
        "introduction": "The first accreting millisecond X-ray pulsar (hereafter AMP) \\sax\\ was discovered in 1998 \\citep{WvdK98, CM98}. Since then, 12 objects of this class have been discovered, with the spin frequencies in the range 182--599 Hz. Many AMPs (e.g., IGR J00291+5934, XTE J1751$-$205, XTE J1807$-$294, \\sax, XTE J1814$-$338) often show nearly sinusoidal profiles, consistent with only one visible  hotspot, probably because the accretion disk extends very close to the stellar surface. The general stability of the pulse profile allows to obtain a high photon statistics and to use the average pulse profile to get constraints on the neutron star mass-radius relation and the equation of state \\citep{PG03}.  However, sometimes the profiles do show variations. In \\sax, for example, a dip appears around the pulse maximum at high photon fluxes, and at low fluxes the profile becomes skewed to the left and sometimes   is double-peaked. The pulse profile evolution  during its five observed outbursts is amazingly similar, with significant changes in the profiles accompanied by large (0.2 cycles) jumps  in the phase of the fundamental (see \\citealt{BDM06,HPC08,HPC09,PW09}; \\citealt{IP09}, hereafter IP09).  Besides these noticeable profile changes, the timing noise with timescales from hours to weeks is known to be present in the data  of this and other AMPs, when the measured pulse phase delays show fluctuations around a mean trend  \\citep{pap07,CCHY08,RDB08,PWvdK09}. There are many possible reasons for the pulse profile changes and the timing noise: changes in the spot shape and position, variations of the emission pattern, inner disk radius or the optical depth of the accretion stream \\citep{p08}. Changes in the  position of the hotspots \\citep{lamb08,PWvdK09} can cause random phase irregularities;  however, it is difficult to imagine that they can lead to qualitative  changes in the pulse form  (such as transformation of a single-peaked profile to a double-peaked)  unless accompanied by strong variation in the  emissivity pattern.  On the other hand, the visibility of the secondary spot  might change even with small fluctuations of the accretion rate. IP09 proposed that the strong pulse profile variations in the 2002 outburst of \\sax\\ results from the appearance of the secondary antipodal spot to our view when the magnetospheric radius increases and the accretion disk recedes from the neutron star with the dropping accretion rate. The decreasing amplitude of Compton reflection in the spectra during the outburst supports this interpretation.  In this Letter, we directly fit the pulse profiles of \\sax\\ during its 2002 outburst with the theoretical model for AMP pulses including the effect of partial screening of the antipodal emitting spot by the accretion disk. We aim at  quantifying the variations in the position of the hotspot centroid, its emissivity pattern and the inner accretion disk radius, and determining the causes of the variations in the pulse shape and the  timing noise.  \\begin{figure} \\centerline{\\epsfig{file= f1.eps,width=8cm}} \\caption{Pulse profiles produced by two antipodal spots of various shapes for $\\rin$ varying from 20 to 30 km with step 2.5 km (from bottom to top). The left hand side panels correspond to the blackbody emissivity pattern $h=0$, and the panels at the right are for $h=0.8$. At lowest radii, the view of antipodal spot is almost fully blocked by the disk, while at highest radii, the whole spot is visible. Stellar parameters: $M=1.4M_\\odot$, $R=10.3$ km, $i=65\\degr$, $\\theta=10\\degr$. The spot angular size $\\rho$ is given by equation (\\ref{eq:rhorin}), and for the ring and crescent geometries,  the inner radius is $\\rho/\\sqrt{2}$.  } \\label{f:profile} \\end{figure} ",
        "conclusions": "We have developed a model for the pulse profile formation in AMPs which accounts for the partial screening of the antipodal emitting spot by the accretion disk. We have demonstrated that the appearance of the antipodal spot, due to the increasing inner disk radius, leads to sharp changes in the pulse profile and corresponding jump in Fourier phases. We showed that a strong secondary maximum appears only in a rather narrow intervals of inner disk radii, explaining the very short appearance of the double-peaked profiles in  \\sax.  By directly fitting the model to the pulse shapes of \\sax\\ observed during its 2002 outburst, we were able to quantify the variations in the position of the spot centroids, their emissivity pattern and the inner accretion disk radius. The sharp jumps in Fourier phases, observed together with the strong changes in the pulse form,  can be explained by the variations of the visibility of the antipodal spot associated with the receding accretion disk.  We also note that a factor of 2 difference in fluxes in 2002 and 2008 outbursts of \\sax, when the pulse profile started to change significantly, implies that the inner disk radius is not a simple function of the accretion rate, but might depend on the previous history. We finally note that  many physical  timing noise mechanisms probably operate in AMPs simultaneously."
    },
    "0910/0910.0160_arXiv.txt": {
        "abstract": "s{I know better than to come between the experts here assembled and their research programs, so I confine these remarks to lessons to be drawn on the state of our subject from the histories of research in three Windows on the Universe:  cosmology, our extragalactic neighborhood, and life in other worlds.} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.5603_arXiv.txt": {
        "abstract": "{ The \\emph{INTEGRAL} satellite is discovering a large population of new X-ray sources which were missed by previous missions due to high obscuration and, in some cases, very short duty cycles. The nature of these sources must be addressed via the characterization of their optical and/or infrared counterparts. } {We investigate the nature of the optical counterparts to five of these newly discovered X-ray sources.} {We combine infrared spectra in the $I$, $J,H$ and $K$ bands together with $JHK$ photometry to characterize the spectral type, luminosity class and distance to the infrared counterparts to these systems. For IGR~J19140$+$0951, we present spectroscopy from the red to the $K$ band and new red and infrared photometry. For SAX~J18186$-$1703 and IGR~J18483$-$0311, we present the first intermediate-resolution spectroscopy reported. Finally, for IGR~J18027$-$2016, we present new $I$ and $K$ band spectra.} {We find that four systems harbour early-type B supergiants. All of them are heavily obscured, with $E(B-V)$ ranging between 3 and 5, implying visual extinctions of $\\sim$ 9 to 15 magnitudes. We refine the published classifications of IGR~J18027$-$2016 and IGR~J19140$+$0951 by constraining their luminosity class. In the first case, we confirm the supergiant nature and rule out class III. In the second case, we propose a slightly higher luminosity class (Ia instead of Iab) and give an improved value of the distance based on new optical photometry. Two other systems, SAX~J18186$-$1703 and IGR~J18483$-$0311 are classified as Supergiant Fast X-ray Transients (SFXTs). XTE~J1901$+$014, on the other hand, contains no bright infrared source in its error circle.} {Owing to their infrared and X-ray characteristics, IGR~J18027$-$2016 and IGR~J19140$+$0951, emerge as Supergiant X-ray binaries with X-ray luminosities of the order of $L_{X}\\sim [1-2]\\times 10^{36}$ erg s$^{-1}$, while SAX~J1818.6$-$1703 and IGR~J18483$-$0311, turn out to be SFXTs at 2 and 3 kpc, respectively. Finally, XTE~J1901$+$014 emerges as a puzzling source: its X-ray behaviour is strongly reminiscent of the SFXTs but a supergiant nature is firmly ruled out for the counterpart. We discuss several alternative scenarios to explain its behaviour. } ",
        "introduction": "The {\\it INTEGRAL} satellite has broadened our knowledge of X-ray binaries, by discovering a significant population of new X-ray sources.  The Third {\\it IBIS/ISGRI} cataloge \\citep{bird07} contains more than 420 sources of which 167 are new.  The on-line cataloge of {\\it INTEGRAL} sources\\footnote{\\tiny{\\texttt{http://isdcul3.unige.ch/$\\sim$rodrigue/html/igrsources.html}}} lists 245 sources discovered or re-discovered by {\\it INTEGRAL}. Of these, 61 ($\\sim 25$ \\%) remain unidentified and 81 ($\\sim 33$ \\%) are extragalactic objects (AGNs, Seyferts, QSOs, etc).  The rest are mainly X-ray binaries. 51 of them ($\\sim 21$ \\%) turn out to be High Mass X-ray Binary candidates (HMXB), which were missed by previous missions because of their high obscuration, because of their very short X-ray cycles or a combination of both. Approximately half of them have been confirmed so far. Amongst these new 51 HMXB systems, there are a few Be/X-ray binaries (BeXB), Supergiant X-ray binaries (SGXB) and the newly established class of Supergiant Fast X-ray Transients \\citep[SFXT; e.g.,][]{smith06,neg06a,sguera06}. This last class comprises 16 candidates of which nearly half have been confirmed.  Before we can conclusively classify every new source and understand its X-ray behavior, we must first establish the nature of the counterpart.   Since the majority of these sources are located in the galactic plane, and concentrated towards the galactic center region, they suffer from large extinction ($A_{V}\\lesssim 15$ mag). Thus, their detection in the visual is often very difficult, and other strategies must be adopted. Detection and study of the counterparts in the infrared, however, is possible with 4~m class telescopes. An important caveat, though, is that the spectral classification of hot stars using infrared spectra is much more uncertain than using the well established MK standard 3950--4750 \\AA\\ region. For example, the spectral classification based on the $K$ band spectrum cannot be done without fundamental ambiguities because of the lack of adequate spectral features in that range \\citep{hanson96}. This problem can be, at least partially, circumvented by using the combined information of several spectral bands. However, this is often impossible to achieve within a single observing campaign.  Following the discovery by \\emph{INTEGRAL}, intense observing campaigns have been undertaken \\citep[amongst others]{chaty08,masetti08,nespoli08,neg09}. As these discoveries proceed, it is fundamental to refine and establish the nature of these new sources. In particular, a very important parameter is the distance to the system which, in turn, gives the X-ray luminosity, the main observable used to constrain the theoretical models.  Most of the work done, however, has been based on $H$ and/or $K$ band spectra to overcome the high extinction. Classification based on a single infrared spectrum (in many cases of intermediate or low resolution only) allows only a crude approximation. For example, the difference in $M_{V}$ between a B1\\,Ib and a B1\\,Ia supergiant (which are difficult or even impossible to tell apart based on an intermediate resolution $H$ or $K$ spectrum) results in an uncertainty of an order of magnitude error in $L_{X}$ \\citep{neg09}. In this paper, we combine new spectroscopy in several bands with new $JHK$ photometry (also $RI$ for one source) to further refine the spectral classification of a sample of these sources. ",
        "conclusions": "\\begin{table*} \\caption[]{Derived parameters for stars in our sample, calculated using the observed data.  } \\label{tab:data} \\begin{center} \\begin{tabular}{llccccccrr} \\hline \\hline \\noalign{\\smallskip} Source & Spectral type & Class & $E(B-V)$ & $E(J-K)$ & $M_{V}$ & $M_{K}$ & $d$ (kpc) & $L_{X}$ (erg/s)$^{\\rm (a)}$ \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} IGR~J18027$-$2016 & B1\\,Ib & SGXB & 3.04$\\pm0.02$ & 1.54 & $-5.8$ & $-5.1$  & 12.4$\\pm 0.1$ &  (1.65$\\pm0.03$) $\\times 10^{36}$ \\\\ SAX~J18186$-$1703 & B0.5\\,Iab & SFXT & 5.08$\\pm0.05$ & 2.54 & $-6.4$ & $-5.65$ & 2.1$\\pm0.1$ & (8$\\pm0.4$) $\\times 10^{35}$ \\\\ IGR~J18483$-$0311 & B0-1\\,Iab & SFXT & 5.22$\\pm0.02$& 2.61 & $-6.4$ & $-5.66$ & 2.83$\\pm 0.05$ & (1.9$\\pm0.06$) $\\times 10^{36}$ \\\\ IGR~J19140$+$0951 & B0.5\\,Iab-a & SGXB & 5.5$\\pm0.1$ & 2.75 & -6.65$\\pm 0.03$ &  -5.90$\\pm0.25$ & 3.6$\\pm 0.04$ & (1.6$\\pm0.3$) $\\times 10^{36}$ \\\\ \\hline \\noalign{\\smallskip} XTE~1901$+$014 & G5\\,III - A3\\,V & LMXB? & $\\sim$3.7 & 1.8 &  &   & 2-7 & 1$^{+1}_{-0.8}\\times 10^{38}$  \\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\begin{list}{}{} \\item[$^{\\mathrm{a}}$] Observed peak luminositites for the energy ranges (2--10)\\,keV, (22--55)\\,keV,  (20--100)\\,keV, (2--20)\\,keV and (3--12) \\,keV respectively.  \\end{list} \\end{center} \\end{table*}   Four of our sources have early-B supergiant companions. All four counterparts turn out to cover a rather narrow spectral interval, being early-B (B0--B1.5) supergiants of moderate or high luminosity. This identifies the nature of the X-ray source unambiguously as HMXBs. Furthermore, they are all heavily obscured, with $E(B-V)\\sim 3-5.5$, implying extinctions on the order of $A_{V}=9-16$ mag in the visual band. But, on the other hand, XTE~1901+014 has no obvious counterpart in any band except $K^{\\prime}$. Thus, a supergiant nature is definitely excluded, thereby setting it apart from the other four.   IGR~J18027$-$2017 is a persistent X-ray source with a counterpart around B1\\,Ib. It is therefore, a SGXB system. The reason why it has not been detected by earlier X-ray missions is not clear. For example, it was not detected by \\emph{ROSAT}. At a distance of $d\\approx 12.4$ kpc, its {\\it XMM-Newton} detected flux in the 2--10~keV band for the main pulse of $8.9\\times 10^{-11}\\:{\\rm erg}\\,{\\rm s}^{-1}\\,{\\rm cm}^{-2}$, would yield a luminosity of $L_{X}=1.65\\times 10^{36}\\:{\\rm erg}\\,{\\rm s}^{-1}$. This is typical for these kind of systems albeit in the lower end. Since the detected flux outside the primary pulse is a little bit smaller, the system can spend a fraction of its time in the upper $\\times 10^{35}$ erg s$^{-1}$. Given its high obscuration, this can explain why it was missed by earlier surveys.    For SAX~J18186$-$1703, the NIR counterpart is a B0.5Iab star, confirming its nature as a SFXT. The 22--50\\,keV flux reported by \\citet{zurita09} is of the order of (2--8)$\\times 10^{-11}\\:{\\rm erg}\\,{\\rm s}^{-1}\\,{\\rm cms}^{-2}$ in quiescence while it reaches (1--2)$\\times 10^{-9}\\:{\\rm erg}\\,{\\rm s}^{-1}\\,{\\rm cms}^{-2}$ at the strongest flares. At the deduced distance of $\\sim 2.1$ kpc, the X-ray luminosity of this object would be $\\sim 3\\times 10^{34}\\:{\\rm erg}\\,{\\rm s}^{-1}$ in quiescence, while the peak luminosity can be as high as $\\sim 8\\times 10^{35}\\:{\\rm erg}\\,{\\rm s}^{-1}$. This is lower than those found in other SFXTs ($\\sim 10^{36}\\:{\\rm erg}\\,{\\rm s}^{-1}$). \\citet{zurita09} have found that this system is in a very eccentric orbit ($e\\sim 0.3-0.4$), with $P_{orb}\\approx 30$ d, amongst the largest yet found for SFXT, reaching a periastron distance between 2 and 3 $R_{*}$. This would locate the compact object at a slightly larger distance from the donor than in the SGXB systems ($a\\lesssim 2R_{*}$) thereby reducing the $L_{X}$ of the outbursts slightly below $10^{36}\\:{\\rm erg}\\,{\\rm s}^{-1}$ as is observed.    IGR~J18483$-$0311 is a transient system with a B0--B1\\,Iab primary. It has been classified as an intermediate SFXT \\citep{rahoui08}. Like many other SFXTs, it is a nearby source at $\\sim$ 2.8 kpc. For this distance, the {\\it INTEGRAL}/IBIS 20--100~keV flux \\citep{sguera07} translates into a luminosity  $\\sim 1.9 \\times 10^{36}\\:{\\rm erg}\\,{\\rm s}^{-1}$, for the strongest flares, which is typical for this kind of objects. The orbital period of this system is 18.5 d, while the spectral type of the companion is similar to that of SAX~J18186$-$1703. Therefore, the compact object would be located slighly closer to the star and, correspondingly, its average X-ray luminosity will be higher, in agreement with the scenario described in \\citet{neg08}.    IGR~J19140$+$0951 is a persistent X-ray source and therefore, owing to its counterpart, a SGXB located at $d\\approx 3-4$~kpc, a little bit further than previously thought. For the bright {\\it INTEGRAL} flux in the 2-2-0~keV band of $1\\times 10^{-9}\\:{\\rm erg}\\,{\\rm s}^{-1}\\,{\\rm cms}^{-2}$ \\citep{hannika04}, the $L_{X}=1-2\\times 10^{36}\\:{\\rm erg}\\,{\\rm s}^{-1}$, typical of a SGXB albeit, again, in the lower end.  The fact that, together with IGR 18027-2017, they have lower X-ray fluxes than the average SGXBs explains their more recent detections in the newer, and deeper {\\it INTEGRAL} observations.  As can be seen in Fig.~\\ref{fig:galaxy}, the sources tend to concentrate in or behind the section of the Sagittarius arm projected onto the direction of the Galactic Bulge, an area intensively scanned by {\\it INTEGRAL}. They are, thus, relatively close to the Sun (in comparison to typical distances to HMXBs), although highly reddened. This fact can introduce a bias in the newly detected populations of obscured sources, which tend to harbor SG companions, as only the brighter and/or closer ones can be reached with telescopes of moderate apertures in the 4~m class. A transient system with a main sequence companion, like a BeXB, will be too faint to be conclusively classified \\citep[i.e.][]{zurita08}. Remarkably, the furthest source in our sample is the least reddened one. This raises the question of the true abundance of HMXBs (and OB stars in fact) in our Galaxy. As can be seen in Table \\ref{tab:data}, the high reddening is not a sufficient condition to be an SFXT since some SGXBs (like IGR~J19140$+$0951) are more reddened. On the other hand, it is not a \\emph {necessary} condition either, as other SFXTs present rather low reddenings \\citep[for example IGR~J08408$-$4503 with $E(B-V)=0.5$][]{neg09}. Furthermore, the position in the arms of the Galaxy of the newly discovered HMXBs (including the SFXTs) is entirely consistent with the positions of the previously known HMXBs (including SGXBs). Therefore there is no reason to believe that they form a different population. This would be expected from a scenario like that discussed in \\citet{neg08}. In this scenario, the SFXTs are explained as just SGXBs where the compact object is located further out in the system, in a region where the stellar wind of the supergiant companion has become substantially clumped. The number of HMXBs in the Galaxy would then be much larger than previously thought. But so would the number of isolated OB stars whose much lower X-ray emission will not show up in the current X-ray surveys if they are heavily obscured. This is supported by the discovery, in the last few years, of a growing population of very massive infrared clusters \\citep[e.g.][]{crow06, mess08}.  Therefore, in principle, there is no reason to believe that the relative abundance of HMXBs, with respect to OB stars, differs from the current values which are well explained by the theoretical models.     Finally, XTE~J1901$+$014 emerges as a transient source whose X-ray behaviour is characteristic of the SFXTs but where the presence of a Supergiant companion is clearly ruled out. It can still be a massive star close to the main sequence in the outskirts of the Galaxy, perhaps with a $16\\,M_{\\sun}$ BH as a companion. In this respect, it is interesting to note the case of the LMC transient A\\,0535$-$66 reaching  $L_{X}\\approx 8\\times 10^{38}\\:{\\rm erg}\\,{\\rm s}^{-1}$ \\citep[e.g.,][]{pp81}. A similar source at the edge of the Galaxy could have a similar hard X-ray behaviour. However, taking into account the moderate hardness (the source is seen as 1RXH~J190140.1+012630) and obscuration of the source, it seems more likely that the counterpart is a late type star located between 2~kpc and the Long Galactic bar (7~kpc), somewhere between luminosity class V to class III, but definitely not consistent with a class I star. Whether a main sequence or giant late type star wind can produce the instabilities necessary to produce the observed outbursts via accretion onto a compact object remains an issue."
    },
    "0910/0910.2812_arXiv.txt": {
        "abstract": "% CCD photometric observations of the Algol-type eclipsing binary X Tri have been obtained. The light curves are analyzed with the Wilson-Devinney (WD) code and new geometric and photometric elements are derived. A new O-C analysis of the system, based on the most reliable timings of minima found in the literature, is presented and apparent period changes are discussed with respect to possible and multiple Light-Time Effect (LITE) in the system. Moreover, the results for the existence of additional bodies around the eclipsing pair, derived from the period study, are compared with those for extra luminosity, derived from the  light curve analysis. ",
        "introduction": "The system was observed during 5 nights from October 2008 to January 2009 at the Athens University Observatory, using a 40-cm Cassegrain telescope equipped with the CCD camera ST-8XMEI and B, V, R, I Bessell filters.  The light curves (hereafter LCs) were analysed with the $PHOEBE~ 0.29d$ software \\citep{PZ05} which uses the 2003 version of the WD code \\citep{WD71,W79}. Due to the lack of spectroscopic mass ratio of the system, the q-search method was applied in Mode 2, 4 and 5 in order to find the most probable value of the (photometric) mass ratio (q).  The least squares method with statistical weights has been used for the analysis of the O-C diagram. The current O-C diagram of X Tri includes 572 times of minima taken from the literature. The following ephemeris: $Min.I=HJD~2422722.285 + 0.9715341^{d} \\times E$ \\citep{K01} was used as the initial one for the O-C analysis of the compiled times of minima. ",
        "conclusions": "The results of the LCs solution (see figure 1 and table 1) show that X Tri is a semi-detached system with the secondary star filling its Roche Lobe. The periodic variations of the orbital period of the system could be explained by adopting the existence of three additional components, which were found to have minimal masses 0.18, 0.24 and 0.22, respectively (see figure 1 and table 2). An extra light contribution to the luminosity of the EB was taken into account in the LCs solution but it was found to be less than 1\\%. Such a small extra luminosity could be explained by the small values of the minimal masses of the possible additional components found. The steady decrease rate of its period is probably due to angular momentum loss, since the direction of the flow (from the more massive to the less massive component), revealed from the O-C diagram analysis, comes in disagreement with the one derived from the LCs analysis.  \\begin{table}  \\caption{The parameters of X Tri derived from the LCs solutions} \\label{tab1} \\scalebox{0.58}{ \\begin{tabular}{ccccccc}  \\hline \\textbf{Parameter}          & \\textbf{value} & \\textbf{Parameter} & \\multicolumn{4}{c}{\\textbf{value}}     \\\\ \\hline $q~(m_{2}/m_{1})$           &    0.599~(2)   &                    &    B    &    V    &    R    &    I     \\\\ $i$~(deg)                   &    87.9~(1)    &   $x_{1}$$^{***}$  &  0.551  &  0.478  &  0.402  &  0.322   \\\\ $T_1$$^{**}$(K), $T_2$ (K)  & 8600, 5188~(4) &   $x_{2}$$^{***}$  &  0.835  &  0.692  &  0.597  &  0.503   \\\\ $A_1$$^*$, $A_2$$^*$        &     1, 0.5     &  $L_{1}/L_{T}$     &0.893~(2)&0.839~(2)&0.795~(2)&0.739~(1) \\\\ $g_1$$^*$, $g_2$$^*$        &     1, 0.32    &  $L_{2}/L_{T}$     &0.107~(1)&0.160~(2)&0.201~(2)&0.246~(3) \\\\ $\\Omega_{1}$, $\\Omega_{2}$  & 4.27~(1), 3.06 &  $L_{3}/L_{T}$     &0.000~(1)&0.000~(1)&0.004~(1)&0.015~(2) \\\\ $\\chi^{2}$                  &     1.278      &                    &                                        \\\\ \\hline \\textit{$^*$assumed},&\\textit{$^{**}$ \\citet{G83},} &\\textit{$^{***}$\\citet{VH93}},& \\textit{$L_{T} = L_{1}+L_{2}+L_{3}$}   \\end{tabular}}  \\end{table}  \\begin{table}  \\caption{The results of the O-C diagram analysis for X Tri} \\label{tab2} \\scalebox{0.52}{  \\begin{tabular}{cccccc}  \\hline \\textbf{Parameters of the EB} & \\textbf{value}  & \\textbf{Parameters of the LITEs} &               \\multicolumn{3}{c}{\\textbf{value}}                \\\\ \\hline $M_1^{*}+M_2~(M_\\odot)$       &  2.1 + 1.26     &                                  &      $3^{rd} body$    &     $4^{th} body$    &    $5^{th} body$ \\\\ $Min. I$~(HJD)                & 2442502.731~(2) &           $P$~(yrs)              &        36.9~(5)       &        22.4~(3)      &       16.8~(4)   \\\\ $P$~(days)                    & 0.9715318~(2)   &   $T_0$~(HJD),~$\\omega$~(deg)    &2452916~(373), 220~(98)&2455069~(335), 34~(13)&     --, --       \\\\ $c_{2}~(\\times 10^{-10})$     &  -2.0308~(2)    &        $A$~(days),~$e$           &  0.0052~(3), 0.2~(2)  &  0.0040~(4), 0.5~(1) &   0.003~(2), 0.0 \\\\ $\\dot{P}~(\\times 10^{-10})$   &  -1.5269~(2)    &      $M_{min}~(M_\\odot)$         &        0.18~(1)       &       0.24~(1)       &      0.22~(1)    \\\\ \\hline \\textit{$^*$assumed}  \\end{tabular}}  \\end{table}   \\begin{figure}[ht!] \\plottwo{LC.eps}{O-C.eps}  \\caption{Left panel: The synthetic LCs (black solid lines) along with the observed ones (colored points). Right panel: The multiperiodic fitting (red solid line) on the O-C points (black points) (upper part) and the O-C residuals (lower part).} \\end{figure}"
    },
    "0910/0910.4025_arXiv.txt": {
        "abstract": "{The majority of stars that leave the main sequence are undergoing extensive mass loss, in particular during the asymptotic giant branch (AGB) phase of evolution. Observations show that the rate at which this phenomenon develops differs highly from source to source, so that the time-integrated mass loss as a function of the initial conditions (mass, metallicity, etc.) and of the stage of evolution is presently not well understood. } {We are investigating the mass loss history of AGB stars by observing the molecular and atomic emissions of their circumstellar envelopes.} {In this work we have selected two stars that are on the thermally pulsing phase of the AGB (TP-AGB) and for which high quality data in the CO rotation lines and in the atomic hydrogen line at 21 cm could be obtained.} {V1942 Sgr, a carbon star of the Irregular variability type, shows a complex CO line profile that may originate from a long-lived wind at a rate of $\\sim$ 10$^{-7}$ \\Msold, and from a young (\\lsim10$^4$\\,years) fast outflow at a rate of $\\sim$ 5 10$^{-7}$ \\Msold. Intense \\HI emission indicates a detached shell with 0.044 \\Msol ~of hydrogen. This shell probably results from the slowing-down, by surrounding matter, of the same long-lived wind observed in CO that has been active during  $\\sim$\\,6\\,10$^{5}$ years. On the other hand, the carbon Mira V CrB is presently undergoing mass loss at a rate of 2\\,10$^{-7}$\\,\\Msold, but was not detected in H\\,{\\sc {i}}. The wind is mostly molecular, and was active for at most 3\\,10$^{4}$ years, with an integrated mass loss of at most 6.5\\,10$^{-3}$\\,\\Msol.} {Although both sources are carbon stars on the TP-AGB, they appear to develop mass loss under very different conditions, and a high rate of mass loss may not imply a high integrated mass loss.} ",
        "introduction": "Low- to intermediate-mass stars, at the end of their main-sequence evolution, become first hydrogen shell-burning red giants (RGB $-$Red Giant Branch$-$ stars), then hydrogen and helium shell-burning red giants (AGB $-$Asymptotic Giant Branch$-$ stars). In this second phase they may undergo mass loss at a very large rate ($>$ 10$^{-8}$ \\Msold), even so large that it has a decisive effect on their late evolution (Olofsson 1999). They are surrounded by expanding envelopes of gas and dust that have been extensively observed with radio molecular lines and infrared continuum emission. These tracers are used to estimate mass-loss rates. However the estimates are somewhat ambiguous because the mass-loss rate of a given source may vary on many different timescales. The mass change as a function of time due to mass loss is thus difficult to evaluate, and to connect with stellar evolution models. Furthermore molecular lines probe an extent of the circumstellar shell (CS) that is limited by photo-dissociation, and therefore furnish information mainly on the inner parts of CSs, and on the recent mass loss.  To try to circumvent these difficulties we have started a systematic programme of observations of red giants in the line of atomic hydrogen at 21 cm. We have published some of our results in several recent papers, and first reports on sizeable samples have been presented by G\\'erard \\& Le~Bertre (2006, hereafter GL2006) and Matthews \\& Reid (2007, hereafter MR2007). A major difficulty of this programme is the confusion caused by the 21 cm emission from the Interstellar Medium (ISM) that is located on the same line-of-sight as the source of interest. This has a strong impact on the observations which have to be conducted with a specific approach, and on the data processing that aims at providing spectra corrected for the ISM emission. Perhaps more confounding, as circumstellar matter is expected to be at some stage injected in the ISM, the confusion by the {\\ub {local}} ISM might actually be at least partly of stellar origin, i.e. caused by material ejected at an earlier stage of evolution.  In addition to observing systematically the \\HI line at 21 cm in a large sample of sources with different properties, it is also important to choose objects for which the Galactic confusion is low and/or can be tracked easily, and therefore corrected accurately. The detailed study of such spectra should serve as a guide to exploit the data that are obtained in more difficult situations.  Here we present our results on two carbon stars, V1942 Sgr and V CrB, for which the confusion is not a serious problem, and which have \\HI properties that differ radically. Both are N-type carbon stars (CGCS 4229 and CCCS 2293, respectively) and have already been detected in CO rotational lines (Olofsson et al. 1993a). However the only available CO spectrum of V1942 Sgr had a poor signal-to-noise ratio, and for our study it appeared essential to also obtain new CO data of high quality. ",
        "conclusions": "The combination of high velocity resolution CO and \\HI data is a promising tool to investigate the history of mass loss by evolved stars. The low level of Galactic \\HI emission and the absence of small-scale structure in this emission have allowed us to obtain \\HI data of high quality on V1942 Sgr and V CrB with the NRT. We have also obtained high quality CO (1-0) and (2-1) spectra of V1942 Sgr with the IRAM 30-m telescope.  For V1942 Sgr, the CO spectra exhibit composite profiles, that reveal a low velocity wind of $\\sim$ 10$^{-7}$ \\Msold ~and a high velocity wind, possibly bipolar, of $\\sim$ 5\\,10$^{-7}$ \\Msold. The comparison with the \\HI spectrum shows that this high velocity wind is recent with an age of at most 10$^4$ years. On the other hand, the low velocity wind appears to have filled a cavity of $\\sim$ 0.2 pc in radius and built the detached shell, that was discovered by IRAS, over a period of 5 10$^5$ years. Follow-up observations with the VLA and ALMA would help to improve this scenario, or possibly lead to a new scheme. The narrowness of the \\HI line profile in V1942\\,Sgr brings new evidence that AGB stellar winds are slowed down by their surrounding medium as surmised by Young et al. (1993b).  For V CrB, the CO spectra that have been published reveal an outflow with expansion velocity, 6.5 \\kms, and mass loss rate, 2.1 10$^{-7}$ \\Msold. The non-detection in \\HI of V CrB sets an upper limit of 3.2 10$^4$ years for the age of this outflow. In such case of a star with low effective temperature, molecular hydrogen data are obviously needed to constrain better the history of mass loss."
    },
    "0910/0910.1436_arXiv.txt": {
        "abstract": "Due to its unique long-term coverage and high photometric precision, observations from the Kepler asteroseismic investigation will provide us with the possibility to sound stellar cycles in a number of solar-type stars with asteroseismology. By comparing these measurements with conventional ground-based chromospheric activity measurements we might be able to increase our understanding of the relation between the chromospheric changes and the changes in the eigenmodes.  In parallel with the Kepler observations we have therefore started a programme at the Nordic Optical Telescope to observe and monitor chromospheric activity in the stars that are most likely to be selected for observations for the whole satellite mission. The ground-based observations presented here can be used both to guide the selection of the special Kepler targets and as the first step in a monitoring programme for stellar cycles. Also, the chromospheric activity measurements obtained from the ground-based observations can be compared with stellar parameters such as ages and rotation in order to improve stellar evolution models. ",
        "introduction": "During recent years there has been a growing interest in understanding stellar cycles in general and the solar cycle in particular. The reason for this is partly the indications that the solar activity cycle could cause climate change (see e.g. Svensmark \\& Friis-Christensen 1997) and partly that solar cycle 24 now seems to be around two years delayed compared to theoretical predictions (Zimmerman 2009).  The ``Sounding stellar cycles with Kepler'' programme will monitor stellar cycles from the amplitude and frequency shifts of the oscillation modes observed in these stars.  By comparing these measurements with conventional ground-based chromospheric activity measurements we might be able to increase our understanding of the relation between the chromospheric changes and the changes in the eigenmodes. Asteroseismic analysis of the Kepler observations also allows us to test the hypothesis put forward by B{\\\"o}hm-Vitense (2007) -- that stellar cycles in active stars are generated by the gradient in the rotation rate close to the surface, whereas the stellar cycles in inactive stars are generated by the gradient in the rotation rate at the base of the convection zone. This can be done as the asteroseismic analysis of the Kepler observations will allow us to measure the depth of the convection zone and give estimates of the internal rotation profile. The details of the ``Sounding stellar cycles with Kepler'' programme are described by Karoff et al. (2009).  The ground-based chromospheric activity measurements can be used not only to compare the changes on the stellar surfaces to those in the interior over the stellar cycles and to test stellar dynamo models, but also to calibrate rotation-activity-age relations. This can be done as asteroseismology offers a unique possibility to measure ages of solar-type field stars to a precision of 5--10\\% (Christensen-Dalsgaard et al. 2007) and because asteroseismology provides the possibility to measure the inclination of the stellar rotation axis so that the spectroscopic measured $v$sin$i$ can be converted into equatorial surface rotation rates (Gizon \\& Solanki 2003). Also, the first ground-based chromospheric activity measurements that we present here on the potential candidates  for observations for the whole mission can be used to rate the different candidates in order to ensure that both active and inactive stars are observed. \\begin{figure}[t] \\begin{center} \\includegraphics[width=3.4in]{karoff.fig01.eps} \\caption{Two spectra illustrating stars with and without chromospheric emission. The wavelength region covers the region around the Ca~{\\sc ii} K line. The upper star (SAO 48081) clearly shows emission in the middle of the Ca~{\\sc ii} K absorption line which proves that this star contains a high level of chromospheric activity. The lower star (SAO 31656) shows no emission which indicates that this star is inactive. The relative intensities have been normalized to the mean intensity in the echelle order and the spectrum for SAO 48081 has been shifted one unit upwards. No radial velocities have been calculated for the stars, but the spectra have been arbitrary aligned.} \\label{fig1} \\end{center} \\end{figure} ",
        "conclusions": ""
    },
    "0910/0910.3605_arXiv.txt": {
        "abstract": "The SXI telescope is one of the three instruments on board EXIST, a multiwavelength observatory in charge of performing a global survey of the sky in hard X-rays searching for Supermassive Black Holes. One of the primary objectives of EXIST is also to study with unprecedented sensitivity the most unknown high energy sources in the Universe, like high redshift GRBs, which will be pointed promptly by the Spacecraft by autonomous trigger based on hard X-ray localization on board. The recent addition of a soft X-ray telescope to the EXIST payload complement, with an effective area of ~950 cm$^2$ in the energy band 0.2-3 keV and extended response up to 10 keV will allow to make broadband studies from 0.1 to 600 keV. In particular, investigations of the spectra components and states of AGNs and monitoring of variability of sources, study of the prompt and afterglow emission of GRBs since the early phases, which will help to constrain the emission models and finally, help the identification of sources in the EXIST hard X-ray survey and the characterization of the transient events detected. SXI will also perform surveys: a scanning survey with sky coverage $\\sim2$~$\\pi$ and limiting flux of $\\sim5\\times10^{-14}$ cgs plus other serendipitous. We give an overview of the SXI scientific performance and also describe the status of its design emphasizing how it has been derived by the scientific requirements. ",
        "introduction": "\\label{sect:sections}  The potential of surveys in high energy astronomy is being fully demonstrated by the sky observations performed at X- and soft $\\Gamma$-ray wavelengths during the last decade. In particular XMM Newton$^1$ with its deep sensitivity in the soft X-ray band, INTEGRAL$^2$ and Swift$^3$ up to $>100$ keV have started to reveal the demographics of sources in the near Universe by studying populations of objects. Still lacking today are both a comprehensive picture of transient phenomena (other than GRBs) spanning a broad range of durations and a deep study of the supermassive objects and their host galaxies against Universe age. Studies of transients require fast repointing and improved sensitivity for proper investigations of their prompt emission. The fast follow-up capability of Swift allows to probe the farthest regions of the Universe by the study of Gamma-Ray Bursts (GRB) and also to discover new populations of transients in our Galaxy. GRBs are today the only beacons that can be used to observe the Universe at $z>3-4$, as demonstrated by Swift e.g. with the latest detection of a GRB at z=8. Swift has opened a new window on the far Universe with GRBs, but many more objects are needed including AGNs.  Swift also carries an X-ray telescope which is an unvaluable tool to observe the GRB afterglows (and Swift has demonstrated that almost all the GRBs have X-ray afterglows).  The Swift/XRT telescope$^4$ is also a powerful tool for localization of the sources seen by Swift/BAT and INTEGRAL/IBIS$^5$ at hard X-ray energies. In the case of IBIS, of the 167 new sources in the 3rd IBIS catalog$^6$ 129 have been identified through optical and NIR spectroscopy. Of these, more than 60\\% have been localized by XRT with typical accuracy $<2-3$ arcsec. Most of the reasons for this success is the possibility for XRT to perform many, relatively short observations.  The concept of operability of Swift is also adopted in the design of the Energetic X-ray Imaging Survey Telescope (EXIST) mission$^7$. EXIST is a proposed hard X-ray imaging all-sky deep survey mission and was recommended by 2001 Report of the Decadal Survey. It is a strong candidate to be the Black Hole Finder Probe, one of the three \"Einstein Probes\" in the Beyond Einstein Program, now proposed for the Astro2010 Decadal Survey. EXIST will be launched in a Low Earth Orbit (LEO) and its primary instrument is the High Energy Telescope (HET), a wide field coded aperture instrument covering the 5-600 keV energy band and imaging sources in a $\\sim70\\times90$deg$^2$ field of view with 2~arcmin resolution and better than 20~arcsec positions$^8$. The energy band of HET overlaps with the soft X-ray range covered by the proposed Soft X-ray Imager (SXI), 0.1-10 keV with an effective area of 950cm$^2$ at 1.5 keV and 3.5m focal length. At longer wavelengths, the IRT$^9$ is an optical-IR aperture telescope covering the $0.3-2.2$ micron range with variable spectral resolution and high sensitivity (AB=24 in 100s). The IRT pixel size is 0.15~arcsec and its Field-of-View in Imaging mode is 16~arcmin$^2$. So EXIST is a real multiwavelength observatory for observations of GRBs and Supermassive Black Holes (SMXB) as well as of many other types of transients and high energy sources.  In the following we will address some of the topics that will be contributed by the SXI telescope in the study of the high energy sources seen by HET and its capabilities for sky survey at soft X-ray wavelengths. Section 2 contains a brief description of the SXI telescope and its effective area, Section 3 and 4 will describe some of the science cases relevant to SXI.  \\begin{figure} \\begin{center} \\begin{tabular}{c} \\includegraphics[height=11cm]{Immagine1_a.eps} \\end{tabular} \\end{center} \\caption[example] { \\label{fig:sxi_telescope} Design overview of the SXI telescope. The instrument (top) has a mirror focal length of 3.5m and the overall length is 4.5m, 70 cm max width. The primary mirror system and the camera (bottom) are also shown. } \\end{figure} ",
        "conclusions": "\\label{sect:sections}  The SXI telescope on board EXIST will take full advantage of the operational strategy adopted for the mission, mostly based on surveys and fast follow-up of GRB and transients. With SXI, EXIST is a real multiwavelength observatory with a sensitive broadband coverage at high energies: 0.1-600 keV,  with SXI providing an effective area of 950cm$^2$ at 1.5keV. SXI will perform wide area surveys (scanning \\& serendipitous) and sensitive observation of transient events in the X-ray (e.g. GRB afterglows, AGN flares). It will also help the identification of HET sources detected during the survey, study the absorbed (even Compton thick) AGN Universe and zoom on AGN states.  The heritage of the Swift/XRT, XMM-Newton and INTEGRAL allows to conclude that the SXI performance is appropriate to the profile of the EXIST mission as currently designed. Main improvements to Swift/XRT are: factor $\\sim10$ effective area, fast detector readout allowing operation during scanning survey."
    },
    "0910/0910.4251.txt": {
        "abstract": "{The growth processes from protoplanetary dust to planetesimals are not fully understood. Laboratory experiments and theoretical models have shown that collisions among the dust aggregates can lead to sticking, bouncing, and fragmentation. However, no systematic study on the collisional outcome of protoplanetary dust has been performed so far so that a physical model of the dust evolution in protoplanetary disks is still missing.} {We intend to map the parameter space for the collisional interaction of arbitrarily porous dust aggregates. This parameter space encompasses the dust-aggregate masses, their porosities and the collision velocity. With such a complete mapping of the collisional outcomes of protoplanetary dust aggregates, it will be possible to follow the collisional evolution of dust in a protoplanetary disk environment.} {We use literature data, perform own laboratory experiments, and apply simple physical models to get a complete picture of the collisional interaction of protoplanetary dust aggregates.} {In our study, we found four different kinds of sticking, two kinds of bouncing, and three kinds of fragmentation as possible outcomes in collisions among protoplanetary dust aggregates. Our best collision model distinguishes between porous and compact dust. We also differentiate between collisions among similar-sized and different-sized bodies. All in all, eight combinations of porosity and mass ratio can be discerned. For each of these cases, we present a complete collision model for dust-aggregate masses between $10^{-12}$ and $10^2$~g and collision velocities in the range $10^{-4} \\ldots 10^4~\\rm cm~s^{-1}$ for arbitrary porosities. This model comprises the collisional outcome, the mass(es) of the resulting aggregate(s) and their porosities.} {We present the first complete collision model for protoplanetary dust. This collision model can be used for the determination of the dust-growth rate in protoplanetary disks.} ",
        "introduction": "} The first stage of protoplanetary growth has still not been fully understood. Although our empirical knowledge on the collisional properties of dust aggregates has considerably widened over the past years \\citep{BlumWurm:2008}, there is no self-consistent model for the growth of macroscopic dust aggregates in protoplanetary disks (PPDs). A reason for such a lack of understanding is the complexity in the collisional physics of dust aggregates. Earlier assumptions of perfect sticking have been experimentally proven false for most of the size and velocity ranges under consideration. Recent work also showed that fragmentation and porosity play important roles in mutual collisions between protoplanetary dust aggregates. In their review paper, \\citet{BlumWurm:2008} show the complex diversity that is inherent to the collisional interaction of dust aggregates consisting of micrometer-sized (silicate) particles. This complexity is the reason why the outcome of the collisional evolution in PPDs is still unclear and why no `grand' theory on the formation of planetesimals, based on firm physical principles, has so far been developed.  The theoretical understanding of the physics of dust aggregate collisions has seen major progress in recent decades. The behavior of aggregate collisions at low collisional energies -- where the aggregates show a fractal nature -- is theoretically described by molecular dynamics simulations of \\citet{DominikTielens:1997}. The predictions of this model -- concerning aggregate sticking, compaction, and catastrophic disruption -- could be quantitatively confirmed by laboratory collision experiments of \\citet{BlumWurm:2000}. Also, the collision behavior of macroscopic dust aggregates was successfully modeled by a smooth particle hydrodynamics method, calibrated by laboratory experiments \\citep{GuettlerEtal:2009a, GeretshauserEtal:preprint}. These simulations were able to reproduce bouncing collisions, which were observed in many laboratory experiments \\citep{BlumWurm:2008}.  However, as laboratory experiments have shown, collisions between dust aggregates at intermediate energies and sizes are characterized by a plethora of outcomes: ranging from (partial) sticking, bouncing, mass transfer, to catastrophic fragmentation \\citep[see][]{BlumWurm:2008}. From this complexity, it is clear that the construction of a simple theoretical model that agrees with all these observational constraints is very challenging. However, in order to understand the formation of planetesimals, it is imperative to describe the entire phase-space of interest, i.e., to consider a wide range of aggregate masses, aggregate porosities, and collision velocities. Likewise, the collisional outcome is a key ingredient of any model that computes the time evolution of the dust size distribution. These collisional outcomes are mainly determined by the collision velocities of the dust aggregates, and these depend on the disk model, i.e. the gas and material density in the disk and the degree of turbulence. Thus, the choice of the disk model (including its evolution) is another major ingredient for dust evolution models.  These concerns lay behind the approach we adopt in this and subsequent papers. That is, instead of first `funneling' the experimental results through a (perhaps ill-conceived) theoretical collision model and then to calculate the collisional evolution, we will directly use the experimental results as input for the collisional evolution model. The drawback of such an approach is of course that experiments on dust aggregate collisions do not cover the whole parameter space and therefore need to be extrapolated by orders of magnitude, based on simple physical models which accuracy might be challenged. However, we feel that this drawback is more than justified by the prospects that our new approach will provide: through a direct mapping of the laboratory experiments, collisional evolution models can increase enormously in their level of realism.  In Paper I, we will classify all existing dust-aggregate collision experiments for silicate dust, including three additional original experiments not published before, according to the above parameters (Sect. \\ref{sec:exp-review}). We will show that we have to distinguish between nine different kinds of collisional outcomes, which we physically describe in Sect. \\ref{sec:exp_types}. For the later use in a growth model, we will sort these into a mass-velocity parameter space and find that we have to distinguish between eight regimes of porous and compact dust-aggregate projectiles and targets. We will present our collision model in Sect. \\ref{sec:collision_regimes} and the consequences for the porosities of the dust aggregates in Sect. \\ref{sec:porosities}. In Sect. \\ref{sec:conclusion}, we conclude our work and give a critical review on our model and the involved necessary simplifications and extrapolations.  In Paper II \\citep{ZsomEtal:2009}, we will then, based upon the results presented here, follow the dust evolution using a recently invented Monte-Carlo approach \\citep{ZsomDullemond:2008} for three different disk models. This is the first fully self-consistent growth simulation for PPDs. The results presented in Paper II represent the state-of-the-art modeling and will give us important insight into questions, such as if the meter-size barrier can be overcome and what the maximum dust-aggregate size in PPDs is, i.e. whether pebbles, boulders, or planetesimals can be formed. ",
        "conclusions": "} In the previous sections we have developed a comprehensive model for the collisional interaction between protoplanetary dust aggregates. The culmination of this effort is Fig. \\ref{fig:colored_regimes}, which presents a general collision model based on 19 different dust-collision experiments, which will be adopted in Paper II. Since it plays a vital role, it is worth a critical appraisal. In a few examples, we want to discuss the main simplifications and shortcomings of our current model.  (1) The categorization into collisions between similar-sized and different-sized dust aggregates (see Figs. \\ref{fig:categorization} and \\ref{fig:colored_regimes}) is well-motivated as we pointed out in Sect. \\ref{sec:collision_regimes}. However, we may ask ourselves whether this binarization is fundamentally correct, if we need more than two categories, or `soft' transitions between the regimes. At this stage, a more complex treatment would be impractical due to the lack of experiments treating this problem.  (2) The binary treatment of porosity (i.e. $\\phi < \\phi_\\mathrm{c}$ for `porous' and $\\phi \\geq \\phi_\\mathrm{c}$ for `compact' dust aggregates) is also a questionable assumption. Although we see fundamental differences in the collision behavior when we use, e.g., porous or compact targets, there might be a smooth transition from the more `porous' to the more `compact' collisions. In addition to that, the assumed value $\\phi_\\mathrm{c} = 0.4$ is reasonable but not empirically affirmed. On top of that, the maximum compaction that a dust aggregate can achieve in a collision depends on many parameters, such as, e.g., the size distribution of the monomer grains \\citep{BlumEtal:2006} and the ability of the granular material to creep sideways inside a dust aggregate \\citep{GuettlerEtal:2009a}.  (3) Although the total number of experiments, upon which our model is based, is unsurpassedly large, the total coverage of parameter space (see the experiment boxes in Fig. \\ref{fig:colored_regimes}) is still small. Thus, we sometimes apply extrapolations into extremely remote parameter-space regions. Although not quantifiable, it must be clear that the error of each extrapolation grows with the distance to the experimentally confirmed domains (i.e. the boxes in Fig. \\ref{fig:colored_regimes}). Clearly, more experiments are required to fill the parameter space, and the identification of the key regions in the mass-velocity plane is exactly one of the goals of Paper II.  (4) With such new experiments, performed at the `hot spots' predicted in Paper II, we will not only close gaps in our knowledge of the collision physics of dust aggregates but will most certainly reveal completely new effects. The rather simple \\cc\\ panel in Fig. \\ref{fig:colored_regimes} as compared to the more complex \\pC\\ is due to the fact that there are hardly any experiments that back-up the \\cc\\ regime, whereas in the \\pC\\ case we have a rather good experimental coverage of the parameter space.  In summary, the sophisticated nature of our collision model is both its strength and its weakness. The drawbacks of identifying four parameters that shape the collision outcome are that rather crude approximations and extrapolations have to be made. However, to acknowledge the role of, e.g., porosity through a binary treatment is still better than to not treat this parameter at all. Our new collision model represents the first attempt to include all existing laboratory experiments (for the material properties of interest); collisional evolution models can enormously profit from this effort.  \\subsection{The Bottleneck for Protoplanetary Dust Growth\\label{sec:outlook}} In this paper, we have presented the framework and physical background for an extended growth simulation. What is to be expected from this? Here, we can speculate under which conditions growth in PPDs is most favorable. A view on Fig. \\ref{fig:colored_regimes} immediately shows that large dust aggregates can preferentially grow for realistic collision velocities in the \\cC\\ and \\pC\\ collision regimes (and to a lesser extent in the \\pc\\ case), due to \\Sd. For this to happen, a broad mass distribution of protoplanetary dust must be present. This prerequisite for efficient growth towards planetesimal sizes has also been suggested by \\citet[][see their Fig. 11]{TeiserWurm:2009a}. Agglomeration experiments with micrometer-sized dust grains and a sticking probability of unity (experiments 1 -- 3 in Table \\ref{tab:experiments}) have shown that nature chooses a rather narrow size distribution for the initial fractal growth phase. If this changes when the physical conditions leave no room for growth under quasi-monodisperse conditions, i.e. whether nature is so `adaptive' and `target-oriented' to find out that growth can only proceed with a wide size distribution, will be the subject of Paper II, in which we apply the findings of this paper to a collisional evolution model.  \\subsection{Influence of the Adopted Material Properties\\label{sec:material_influence}} The choice of material in our model is 1.5~$\\mu$m diameter silica dust as most of the underlying experiments were performed with this material. Many experiments \\citep{BlumWurm:2000, LangkowskiEtal:2008, BlumWurm:2008} showed that this material is at least in a qualitative sense representative for other silicatic materials -- also for irregular grains with a broader size distribution. Still, the grain size of the dust material may have a quantitative influence on the collisional outcomes. For example, dust aggregates consisting of 0.1~$\\mu$m are assumed to be stickier and more rigid \\citep{WadaEtal:2007, WadaEtal:2008, WadaEtal:2009}, because the grain size may scale the rolling force or breaking energy entering into Eqs. \\ref{eq:S1_threshold} and \\ref{eq:S2_threshold}. However, due to a lack of experiments with smaller monomer sizes, we cannot give a scaling for our model for smaller monomer sizes at this point. Moreover, organic or icy material in the outer regions of PPDs or oxides and sintered material in the inner regions may have a big impact on the collisional outcome, i.e. in enhancing the stickiness of the material and thereby potentially opening new growth channels.  As for organic materials, \\citet{KouchiEtal:2002} found an enhanced sticking of cm-sized bodies covered with a 1~mm thick layer of organic material at velocities as high as 500~\\cms\\ and a temperature of $\\sim250$~K. Also icy materials are likely believed to have an enhanced sticking efficiency compared to silicatic materials. \\citet{HatzesEtal:1991} collided 5~cm diameter solid ice spheres, which were covered with a 10 -- 100~$\\mu$m thick layer of frost. They found sticking for a velocity of 0.03~\\cms, which is in a regime where our model for refractory silicatic material predicts bouncing (see \\pp\\ or \\cc\\ in Fig. \\ref{fig:colored_regimes}). Sintering of porous dust aggregate may occur in the inner regions near the central star or -- triggered by transient heating events \\citep[e.g. lightning,][]{GuettlerEtal:2008} -- even further out. Ongoing studies with sintered dust aggregates \\citep{Poppe:2003} show an increased material strength (e.g. tensile strength) by an order of magnitude (C. G\\\"uttler \\& J. Blum, unpublished data). This would at least make the material robust against fragmentation processes and qualitatively shift them from the porous to the compact regime in our model -- without necessarily being compact. Due to a severe lack on experimental data for all these materials, it is necessary and justified to restrict our model to silicates at around 1~AU while it is to be kept in mind that these examples of rather unknown materials might potentially favor growth in other regions in PPDs."
    },
    "0910/0910.1093_arXiv.txt": {
        "abstract": "We provide fits to the distribution of galaxy luminosity, size, velocity dispersion and stellar mass as a function of concentration index $C_r$ and morphological type in the SDSS.  (Our size estimate, a simple analog of the SDSS {\\tt cmodel} magnitude, is new:  it is computed using a combination of seeing-corrected quantities in the SDSS database, and is in substantially better agreement with results from more detailed bulge/disk decompositions.) We also quantify how estimates of the fraction of `early' or `late' type galaxies depend on whether the samples were cut in color, concentration or light profile shape, and compare with similar estimates based on morphology. Our fits show that ellipticals account for about 20\\% of the $r$-band luminosity density, $\\rho_{L_r}$, and 25\\% of the stellar mass density, $\\rho_*$; including S0s and Sas increases these numbers to 33\\% and 40\\%, and 50\\% and 60\\%, respectively. The values of $\\rho_{L_r}$ and $\\rho_*$, and the mean sizes, of E, E+S0 and E+S0+Sa samples are within 10\\% of those in the Hyde \\& Bernardi (2009), $C_r\\ge 2.86$ and $C_r\\ge 2.6$ samples, respectively. Summed over all galaxy types, we find $\\rho_* \\sim 3\\times 10^8 M_\\odot$Mpc$^{-3}$ at $z\\sim 0$. This is in good agreement with expectations based on integrating the star formation history. However, compared to most previous work, we find an excess of objects at large masses, up to a factor of $\\sim 10$ at $M_*\\sim 5 \\times 10^{11}M_\\odot$. The stellar mass density further increases at large masses if we assume different IMFs for elliptical and spiral galaxies, as suggested by some recent chemical evolution models, and results in a better agreement with the dynamical mass function.  We also show that the trend for ellipticity to decrease with luminosity is primarily because the E/S0 ratio increases at large $L$.  However, the most massive galaxies, $M_*\\ge 5\\times 10^{11}M_\\odot$, are less concentrated and not as round as expected if one extrapolates from lower $L$, and they are not well-fit by pure deVaucouleur laws.  This suggests formation histories with recent radial mergers. Finally, we show that the age-size relation is flat for ellipticals of fixed dynamical mass, but, at fixed $M_{\\rm dyn}$, S0s and Sas with large sizes tend to be younger.  Hence, samples selected on the basis of color or $C_r$ will yield different scalings.  Explaining this difference between E and S0 formation is a new challenge for models of early-type galaxy formation. ",
        "introduction": "Each galaxy has its own peculiarities. Nevertheless, even to the untrained eye, sufficiently well-resolved galaxies can be separated into three morphological types: disky spirals, bulgy ellipticals, and others which are neither. The morphological classification of galaxies is a field that is nearly one hundred years old, and sample sizes of a few thousand morphologically classified galaxies are now available (e.g. Fukugita et al. 2007; Lintott et al. 2008).  However, such eyeball classifications are prohibitively expensive in the era of large scale sky surveys, which image upwards of a few million galaxies.  Moreover, the morphological classification of even relatively low redshift objects from ground-based data is difficult. Thus, a number of groups have devised automated algorithms for discerning morphologies from such data (e.g. Ball et al. 2004 and references therein).  In parallel, it has been recognized that relatively simple criteria, using either crude measures of the light profile (e.g. Strateva et al. 2001), the colors (e.g. Baldry et al. 2004), or some combination of photometric and spectroscopic information (Bernardi et al. 2003; Bernardi \\& Hyde 2009) allow one to separate early-type galaxies from the rest.  Because they are so simple, these tend to be more widely used. The main goal of this paper is to show how samples based on such crude cuts compare with those which are based on the eyeball morphological classifications of Fukugita et al. (2007). We do so by comparing the luminosity, stellar mass, size and velocity dispersion distributions for cuts based on photometric parameters with those based on morphology.  These were chosen because the luminosity function is standard, although it is becoming increasingly common to compare models with $\\phi(M_*)$ rather than $\\phi(L)$ (e.g. Cole et al. 2001; Bell et al. 2003; Panter et al. 2007; Li \\& White 2009); the size distribution $\\phi(R)$ has also begun to receive considerable attention recently (Shankar et al. 2009b); and the distribution of velocity dispersions $\\phi(\\sigma)$ (Sheth et al. 2003) is useful, amongst other things, to reconstruct the mass distribution of super-massive black holes (e.g. Shankar et al. 2004; Tundo et al. 2007; Bernardi et al. 2007; Shankar, Weinberg \\& Miralda-Escud{\\'e} 2009; Shankar et al. 2009a) and in studies of gravitational lensing (Mitchell et al. 2005).  Section~\\ref{data} describes the dataset, the photometric and spectroscopic parameters derived from it, and the subsample defined by Fukugita et al. (2007). This section shows how we use quantities output from the SDSS database to define seeing-corrected half-light radii which closely approximate the result of bulge + disk decompositions. We describe our stellar mass estimator in this section as well; a detailed comparison of it with stellar mass estimates computed by three different groups is presented in Appendix~\\ref{Mscomp}. The result of classifying objects into two classes, on the basis of color, concentration index, or morphology are compared in Section~\\ref{sec:mcba}. Luminosity, stellar mass, size and velocity dispersion distributions, for the Fukugita et al. morphological types are presented in Section~\\ref{DFs}, where they are compared with those based on the other simpler selection cuts.  This section includes a discussion of the functional form, a generalization of the Schechter function, which we use to fit our measurements. We find more objects with large stellar masses than in previous work (e.g. Cole et al. 2001; Bell et al. 2003; Panter et al. 2007; Li \\& White 2009); this is the subject of Section~\\ref{phiMass}, where implications for the match with the integrated star formation rate, and the question of how the most massive galaxies have evolved since $z\\sim 2$ are discussed.  While we believe these distributions to be interesting in their own right, we also study a specific example in which correlations between quantities, rather than the distributions themselves, depend on morphology.  This is the correlation between the half-light radius of a galaxy and the luminosity weighted age of its stellar population.  Section~\\ref{ageRe} shows that the morphological dependence of this relation means it is sensitive to how the `early-type' sample was selected, potentially resolving a discrepancy in the recent literature (Shankar \\& Bernardi 2009; van der Wel et al. 2009; Shankar et al. 2009c). A final section summarizes our results, many of which are provided in tabular form in Appendix~\\ref{tables}.  Except when stated otherwise, we assume a spatially flat background cosmology dominated by a cosmological constant, with parameters $(\\Omega_m,\\Omega_\\Lambda)=(0.3,0.7)$, and a Hubble constant at the present time of $H_0=70$~km~s$^{-1}$Mpc$^{-1}$. When we assume a different value for $H_0$, we write it as $H_0=100h$~km~s$^{-1}$Mpc$^{-1}$. ",
        "conclusions": "\\label{discuss} We compared samples selected using simple selection algorithms based on available photometric and spectroscopic information with those based on morphological information. Requiring concentration indices $C_r\\ge 2.6$ selects a mix in which E+S0+Sa's account for about two-thirds of the objects; if $C_r\\ge 2.86$ instead, then two-thirds of the sample comes from E+S0s; whereas Es alone account for more than two-thirds of a sample selected following Hyde \\& Bernardi (2009) (Figures~\\ref{fredC} and~\\ref{fblueC}, and Table~\\ref{purity}).  E's alone account for about 40\\%, 50\\% and 75\\% of the total stellar mass in samples selected in these three ways.  The reddest objects at intermediate luminosities or stellar masses are edge-on disks (Figure~\\ref{gmrmorph}).  As a result, samples selected on the basis of color alone, or cuts which run parallel to the red sequence are badly contaminated by such objects. However, simply adding the additional requirement that the axis ratio $b/a\\ge 0.6$ is an easy way to remove such red edge-on disks from the `red' sequence; the resulting sample is similar to requiring $C_r\\ge 2.86$.  This may provide a simple way to select relatively clean early-type samples in higher redshift datasets (e.g. DEEP2, $z$Cosmos). Our measurements provide the low redshift benchmarks against which such future higher redshift measurements can be compared.  We showed how the distribution of luminosity, stellar mass, size and velocity dispersion in the local universe is partitioned up amongst different morphological types, and we compared these distributions with those based on simple selection algorithms based on available photometric and spectroscopic information (Figures~\\ref{differentCIa}--\\ref{ES0Sa}). We described our measurements by assuming that the intrinsic distributions have the form given by equation~(\\ref{phiX}). We showed how measurement errors bias the fitted parameters (equation~\\ref{psiO}), and used this to devise a simple method which removes this bias.  The results, which are reported in tabular form in Appendix~\\ref{tables}, show that ellipticals contain $\\sim$20\\% of the luminosity density and 25\\% of the stellar mass density in the local universe, and have mean sizes of order 3.2~kpc.  Including S0s increases these numbers to 33\\% and 40\\%; adding Sas results in further increases, to 50\\% and 60\\% respectively. These numbers are in broad agreement with those from the Millennium Galaxy Survey of about $10^4$ objects in $37.5$~deg$^2$. Driver et al. (2007) report that $15\\pm 5$\\% of the stellar mass density is in ellipticals, and adding bulges increases this to $44\\pm 9$\\%.   Our stellar mass function has more massive objects than other recent determinations (e.g. Cole et al. 2001; Bell et al. 2003; Panter et al. 2007; Li \\& White 2009), similarly shifted to a Chabrier (2003) IMF (Figures~\\ref{MsFpet} and ~\\ref{MsFmod}). The mass scale on which the discrepancy arises is of order where some previous work had only a handful of objects -- our substantially larger volume is necessary to provide a more reliable estimate of these abundances. Using stellar masses estimated from {\\tt cmodel} luminosities, which are more reliable than Petrosian luminosities at the large masses where the discrepancy in $\\phi(M_*)$ is largest, gives stellar mass densities in objects more massive than $(1,2,3)\\times 10^{11}\\,M_\\odot$ that are larger by more than $\\sim$~$(20,50,100)$ percent compared to Bell et al. (2003) (Figure~\\ref{MsFcumulative}).  This analysis required that we study the sytematic differences between the stellar mass estimates based on $g-r$ color (our equation~\\ref{gmrBell}, following Bell et al. 2003), colors in multiple bands (Blanton \\& Roweis 2007), and on spectral features (Gallazzi et al. 2005). (See Gallazzi \\& Bell 2009, which appeared while our work was being refereed, for a discussion of the pros and cons of these various approaches, and of the accuracy to which stellar masses can currently be derived.) The $g-r$ and Gallazzi et al. estimates are generally in good agreement (Figure~\\ref{BG}), although the spectral based estimates suffer slightly from aperture effects which are complicated by the magnitude limit of the survey (Figures~\\ref{BGz}, \\ref{CompareMsz} and~\\ref{MsFpet}). The Blanton et al. estimates are in good agreement with the other two provided one uses LRG-based templates to estimate masses at the most massive end (Figures~\\ref{BG-BR}, \\ref{CompareMsAll} and~\\ref{MsFmod}). At lower masses, some combination of the LRG and other templates is required.  Ignoring the LRG templates altogether (e.g. Li \\& White 2009) results in systematic underestimates of as much as 0.1~dex or more (Figure~\\ref{BG-BR}), severely compromising estimates of the number of stars currently locked up in massive galaxies (Figure~\\ref{MsFpet}). If we compare our estimate of the stellar mass density in objects more massive than $(1,2,3)\\times 10^{11}\\,M_\\odot$ with those from the Li \\& White (2009) fit, then our values are $\\sim$~$(140,230,400)$ percent larger.  Allowing more high mass objects means that major dry mergers may remain a viable formation mechanism at the high mass end (Figure~\\ref{Bezanson}).  It also relieves the tension between estimates of the evolution of the most massive galaxies which are based on clustering (which predict some merging, and so some increase in stellar mass; Wake et al. 2008; Brown et al. 2008) and those based on abundances (for which comparison of high redshift measurements with the previous $z\\sim 0$ measurements indicated little evolution; Wake et al. 2006; Brown et al. 2007; Cool et al. 2008).  This discrepancy may be related to the origin of intercluster light (e.g. Skibba et al. 2007; Bernardi 2009); our measurement of a larger local abundance in galaxies reduces the amount of stellar mass that must be stored in the ICL.  It has been argued that a number of observations are better reproduced if one assumes a different IMF for elliptical and spiral galaxies (e.g. Calura et al. 2009).  We showed that this acts to further increase the abundance of massive galaxies (Figure~\\ref{MsFIMF}), and reduces the difference between stellar and dynamical mass, especially at larger masses.  At $M_*\\ge 10^{11}M_\\odot$, the increase due to the change in IMF is a factor of two with respect to models which assume a fixed IMF.  If we sum up the observed counts to estimate the stellar mass density in the range  $8.6 < \\log_{10} M_* < 12.2$ $M_{\\odot}$ ($M_*$ from equation~\\ref{gmrBell} using {\\tt cmodel} magnitudes), then the result is $3.05\\times 10^8M_\\odot$~Mpc$^{-3}$. Using our fit to the observed distribution (values between round brackets in Table~\\ref{tabMs}) gives a similar value ($3.06\\times 10^8M_\\odot$~Mpc$^{-3}$) and a slightly smaller value if one uses the {\\it intrinsic} fit ($2.89\\times 10^8M_\\odot$~Mpc$^{-3}$, see Table~\\ref{tabMs}). Our values are $\\sim 15$\\% and 30\\% larger than those reported by Panter et al. (2007) and Li \\& White (2009), respectively. If we allow a type dependent IMF, the total stellar mass density increases by a further 30\\%.  However, although our stellar mass function has more $M_* > 10^{11}$ M$_\\odot$ objects than other recent determinations, our estimate of the total stellar mass density is similar to that measured by Bell et al. (2003). It is about 20\\% smaller than the value reported by Driver et al. (2007) (once shifted to the same IMF, for which we have chosen Chabrier 2003; see Table~\\ref{tabIMF}). This is because differences at the mid/faint end contribute more to the total stellar density than the difference we measured at the massive end.  It has been suggested that direct integration of the cosmological star formation rate overpredicts the total local estimate of the stellar mass density (see, e.g., Wilkins et al. 2008, and references therein). However, we showed that recent determinations of the recycling factor (equation~\\ref{eq|rhoz}) and the high-$z$ star formation rate (Figure~\\ref{fig|SFRz}) result in better agreement (Figure~\\ref{fig|rhoz}). This is because the former yields smaller remaining masses, and the latter produces fewer stars formed in the first place.  Our measurements also show that the most luminous or most massive galaxies, which one might identify with BCGs, are less concentrated and have smaller $b/a$ ratios, than slightly less luminous or massive objects (Figures~\\ref{cmorph}, \\ref{bamorph} and~\\ref{fredC}). Their light profile is also not well represented by a pure deVaucoleur law. This is consistent with results in Bernardi et al. (2008) and Bernardi (2009) who suggest that these are signatures of formation histories with recent radial mergers. In this context, note that we showed how to define seeing-corrected sizes, using quantities output by the SDSS pipeline, that closely approximate deVaucouleur bulge + Exponential disk decompositions (equations~\\ref{Rcmodel} and~\\ref{skysubr}).  Our {\\tt cmodel} sizes represent a substantial improvement over Petrosian sizes (which are not seeing corrected) and pure {\\tt deV} or {\\tt Exp} sizes (Figures~\\ref{cmodel} and~\\ref{cmodel2}).  And finally, our study of the age-size correlation resolves a discrepancy in the literature:  whereas Shankar et al. (2009c) report no correlation at fixed $M_{\\rm dyn}$, van der Wel et al. (2009) report that larger galaxies tend to be younger.  We showed that ellipticals follow the scaling reported by Shankar et al. scaling, whereas S0s and Sas follow that of van der Wel et al. (Figure~\\ref{ageRMd}), suggesting that Shankar et al. select a sample dominated by galaxies with elliptical morphologies, whereas van der Wel et al. include more S0s and Sas.  These conclusions about the differences between the samples are consistent with how the samples were actually selected.  Since van der Wel et al. use their measurements to constrain a model for early-type galaxy formation, this is an instance in which having morphological information matters greatly for the physical interpretation of the data.  Our results indicate that models of early-type galaxy formation should distinguish between ellipticals and S0s because, in the projection of the age-size-mass correlation shown in Figure~\\ref{ageRMd}, the S0s and Sas are very different from the other morphological types.  Whether the smaller sizes for older S0s are due to the gradual stripping away of a younger disk is an open question.  van der Wel et al. (2009) use their observation that the age-size relation at fixed $\\sigma$ is flat to motivate a model in which early-type galaxy formation requires a critical velocity dispersion (which they allow to be redshift dependent).  The same logic applied here suggests that while this may be reasonable for S0s or Sas, it is not well-motivated for ellipticals (Figure~\\ref{ageRsig}). However, it might be interesting to explore a model in which elliptical formation requires a critical (possibly redshift dependent) dynamical mass rather than velocity dispersion (which may be redshift dependent).  This is interesting because, in hierarchical models, the phenomenon known as down-sizing (e.g., Cowie et al. 1996; Heavens et al. 2004; Sheth et al. 2006; Jimenez et al. 2007) is then easily understood (Sheth 2003)."
    },
    "0910/0910.4480_arXiv.txt": {
        "abstract": "A search for muon neutrinos from Kaluza-Klein dark matter annihilations in the Sun has been performed with the 22-string configuration of the IceCube neutrino detector using data collected in 104.3 days of live-time in 2007. No excess over the expected atmospheric background has been observed. Upper limits have been obtained on the annihilation rate of captured lightest Kaluza-Klein particle (LKP) WIMPs in the Sun and converted to limits on the LKP-proton cross-sections for LKP masses in the range 250 -- 3000 GeV. These results are the most stringent limits to date on LKP annihilation in the~Sun. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.4449_arXiv.txt": {
        "abstract": "{} {We investigate the temporal evolution of magnetic flux emerging within a granule in the quiet-Sun internetwork at disk center.} {We combined IR spectropolarimetry of high angular resolution performed in two \\ion{Fe}{i} lines at 1565\\,nm with speckle-reconstructed G-band imaging. We determined the magnetic field parameters by a LTE inversion of the full Stokes vector using the SIR code, and followed their evolution in time. To interpret the observations, we created a geometrical model of a rising loop in 3D. The relevant parameters of the loop were matched to the observations where possible. We then synthesized spectra from the 3D model for a comparison to the observations.} {We found signatures of magnetic flux emergence within a growing granule. In the early phases, a horizontal magnetic field with a distinct linear polarization signal dominated the emerging flux. Later on, two patches of opposite circular polarization signal appeared symmetrically on either side of the linear polarization patch, indicating a small loop-like structure. The mean magnetic flux density of this loop was roughly 450\\,G, with a total magnetic flux of around 3$\\times 10^{17}$\\,Mx. During the $\\sim$12\\,min episode of loop occurrence, the spatial extent of the loop increased from about 1 to 2\\,arcsec. The middle part of the appearing feature was blueshifted during its occurrence, supporting the scenario of an emerging loop. There is also clear evidence for the interaction of one loop footpoint with a preexisting magnetic structure of opposite polarity. The temporal evolution of the observed spectra is reproduced to first order by the spectra derived from the geometrical model. During the phase of clearest visibility of the loop in the observations, the observed and synthetic spectra match quantitatively.} {The observed event can be explained as a case of flux emergence in the shape of a small-scale loop. The fast disappearance of the loop at the end could possibly be due to magnetic reconnection.} ",
        "introduction": "\\label{introduction}  Our knowledge of small-scale magnetic fields located in the photosphere and low chromosphere has evolved extremely fast thanks to the high spatial resolution observations done with the Hinode satellite (e.g., Lites et al. \\cite{litesetal08}), and the improvement of the spatial resolution of large ground-based solar telescopes achieved by adaptive optics systems (e.g., von der L{\\\"u}he et al. \\cite{vdlueheetal03}). The results of such observations tracing the weakest magnetic fields can now be compared directly with theoretical models, for example on the question whether a local dynamo operates in the solar photosphere (Cattaneo \\cite{cattaneo99}). The observations may also help to refine modern numerical simulations of surface magnetoconvection (V\\\"{o}gler et al. \\cite{vogleretal05}, Schaffenberger et al. \\cite{schaffenbergeretal06}, Stein \\& Nordlund \\cite{stein_nordlund06}, Abbett \\cite{abbett07}) by providing more stringent boundary conditions on magnetic field properties in the solar photosphere.  Direct measurements of the 3D topology and the temporal evolution of magnetic loops associated with small granular and intergranular structures are, however, still rare. The measurements require accurate spectropolarimetric 2D data of all Stokes parameters with high spatial resolution, high signal-to-noise ratio, and high cadence. Mart\\'{\\i}nez Gonz\\'{a}lez et al. (\\cite{martinezgonzalezetal07}) reported observations of the magnetic field vector in the internetwork. They showed evidence that small-scale (only 2-6\\,arcsec long) low-lying loops connect at least 10-12\\% of the internetwork magnetic flux measured at 1565\\,nm.  Centeno et al. (\\cite{centenoetal07}, CE07) and Orozco Su\\'{a}rez et al. (\\cite{orozcosuarezetal08}, OR08) analyzed the temporal evolution of such small-scale magnetic features. They used similar data sets of the \\ion{Fe}{i}\\,630.15\\,nm and 630.25\\,nm spectral lines from the Hinode/SOT spectropolarimeter (Lites et al. \\cite{litesetal01}, Kosugi et al. \\cite{kosugietal07}). OR08 % constructed maps of the temporal evolution of linear and circular polarization signals, line-of-sight velocity, and the intensities of the \\ion{Fe}{i}\\,630.25\\,nm and \\ion{Ca}{ii}\\,H spectral line. They found magnetic signals appearing at the central parts of granules, but without a significant linear polarization signal. Estimated lifetimes of the events were 20 and 14\\,min. They did not find any significant coupling of the events to the chromosphere. CE07 % applied an inversion assuming local thermodynamic equilibrium (LTE) to the Hinode spectra to retrieve information on the temporal evolution of the magnetic flux and its topology. Similar to OR08, % they also reported the appearance of magnetic signals in the central parts of granules, and documented a drift of vertical dipoles toward the surrounding intergranular lanes. CE07 found a slightly shorter lifetime of only 8 min. Moreover, they found that the appearance of horizontal magnetic fields precedes that of the vertical fields. They interpreted their results as the rise of small-scale magnetic loops through the photosphere, although they could not exclude a descending loop.  In this paper, we present a similar analysis of the temporal evolution of the appearance and topology changes during a particular flux emergence event in the internetwork at the disk center, using ground-based observations in the infrared spectral region rendering high polarimetric precision and magnetic sensitivity. We compare these observations with synthetic spectra from a geometrical model of a rising loop to investigate whether the observations can be interpreted in such terms. ",
        "conclusions": "\\label{discus_and_concl}  High-resolution surveys of large areas of quiet-Sun internetwork (Lites et al. \\cite{litesetal08}) and plage regions (Ishikawa et al. \\cite{ishikawaetal08}) revealed widespread occurrences of relatively strong horizontal fields over granules, accompanied by vertical bipolar fields cospatial with up- and downflows (Mart\\'{\\i}nez Gonz\\'{a}lez et al. \\cite{martinezgonzalezetal07}). On the other hand, individual granules are often collocated with strong unipolar fields decoupled from granular outflows (OR08, % Bello Gonz\\'{a}lez et al. \\cite{bellogonzalezetal08}). \\begin{figure} \\centerline{\\resizebox{8.8cm}{!}{\\includegraphics{small-12807f12.eps}}} \\caption{Comparison of observed ({\\em bottom}) and simulated spectra ({\\em top}). {\\em Left to right}: Stokes $IQUV$. {\\em Bottom to top}: Time steps 1 to 6. The display ranges for $QUV$ are given at {\\em bottom right} in each panel. The {\\em vertical dashed line} in Stokes $I$ denotes the rest wavelength of the 1564.8\\,nm line. The short {\\em horizontal bars} at the left in $IQUV$ of step 4 denote the location of the profiles shown in Fig.\\,\\ref{fig_profiles}. } \\label{spec_comp} \\end{figure}  CE07 % reported for the first time the emergence of a magnetic loop inside of a granule whose footpoints drifted towards intergranular lanes. Here we present new observational evidence for a similar event, whose onset manifests itself through enhanced net linear polarization within an evolved granule. Soon after, a pair of opposite circularly polarized patches appear symmetrically on each side of the linear polarization forming a small magnetic dipole bridged by a horizontal magnetic field. As long as both patches of circular polarization are visible, they are about of the same size and show a balanced magnetic flux of opposite sign ($+2.9$ and $-3.1\\times 10^{17}$ Mx in step 4). This supports the conclusion made by Lamb et al. (\\cite{lambetal08}) that the model of an asymmetric dipole emergence is very unlikely.  The fast disappearance of the linear polarization signal after 10:43:53\\,UT (step 5, see Figs.\\,\\ref{fig_tip_overview}, \\ref{fig_gband_magparam}, or \\ref{2d_maps}) is in striking contrast with its gradual growth seen from 10:38:47\\,UT to 10:43:53\\,UT. Can we understand it just by an upward extending loop reaching up to a higher atmospheric layer undetectable by the \\ion{Fe}{i} lines as suggested in CE07? % If we consider the loop top as a magnetic perturbation propagating upwards through otherwise non-magnetic layers, then a smooth decay of contribution and response functions with height would imply slower ceasing of the net linear polarization, perhaps detectable on the next frame. Since the contribution and response functions of Stokes {\\em Q} and {\\em U} to magnetic perturbations can be very complex (del Toro Iniesta \\cite{deltoroiniesta_book03}), their detailed computations would help to understand the disappearance of linear polarization after 10:43:53\\,UT. But after 10:43:53\\,UT the circular polarization signal of the footpoint that meets an opposite-polarity magnetic field also weakens rapidly. This could indicate magnetic reconnection between the emergent flux and the preexisting magnetic fields. If reconnection thus is such a common feature, it could explain why OR08 found no trace of the flux emergence in the chromospheric layers.  The disappearance of the polarization signal could, however, also be related to other causes. In the rising loop simulation, the linear polarization amplitude at the last step is again comparable to the first one. Taking into account the noise present in the observations, the polarization signal could also be simply reduced below the detection limit. Another possibility could be a slight spatial drift of the feature with time that would remove it partly out of the observed FoV, because it already was located only in the first three steps of each repeated scan (see Fig.\\,\\ref{fig_tip_overview}). Moreover, the possible expansion of the magnetic loop with height (and time) could also affect the measurement of the linear signal above the detection threshold. From the comparison with the simulated loop, we conclude, however, that at least the ascent until maximal linear polarization signal is reached is fully traced by the observations.  Our flux density estimate of $\\sim$450\\,G found in the emerging loop is comparable to the values reported in Mart\\'{\\i}nez Gonz\\'{a}lez et al. (\\cite{martinezgonzalezetal07}) and Ishikawa et al. (\\cite{ishikawaetal08}) for small-scale magnetic loops in the solar internetwork and plage region, respectively; it significantly exceeds, however, the 10-100\\,G given by Mart\\'{\\i}nez Gonz\\'{a}lez \\& Bellot Rubio (\\cite{martinezgonzalez_bellotrubio09}). This could be due to the assumption of a magnetic filling factor of unity by the latter authors. The loop footpoints separate with time by around 2\\,km\\,s$^{-1}$, drifting in opposite directions towards intergranular lanes. Moreover, our estimate of the magnetic flux density shows that the field is stronger than the typical equipartition value of 400\\,G (Cheung et al. \\cite{cheungetal07}, Ishikawa et al. \\cite{ishikawaetal08}). This suggests that both granular outflow and relaxation of the magnetic tension of the loop may drive the footpoints.  How does our loop compare to the one reported in CE07? % They occupy similar locations and enclose too little magnetic flux to leave any trace on the underlying granules or ambient G-band bright points. On the other hand, our loop displays a higher magnetic flux density and a longer lifetime. \\begin{figure} \\resizebox{8.8cm}{!}{\\includegraphics{small-12807f13.eps}} \\caption{Comparison of observed ({\\em black}) and synthetic spectra ({\\em gray}) in step 4. {\\em Top to bottom}: Stokes $IQUV$. {\\em Left to right}: Lower footpoint, center of loop, upper footpoint. } \\label{fig_profiles} \\end{figure}  Another difference to the loop reported in CE07 % is that in our case one loop footpoint meets the ambient field of opposite polarity, which is followed by the fast disappearance of the loop, in contrast to its more gradual emergence. This is very symptomatic of field cancellation by magnetic reconnection, as reported in Cheung et al (\\cite{cheungetal08}) and Guglielmino et al. (\\cite{guglielminoetal08}). The magnetic reconnection likely drags loop remnants very rapidly out of the line-forming domain less than 12 min after the first trace of the loop top was detected.  The magnetic loop emergence presented here is probably just a particular event in a multitude of similar ones occurring continuously in the solar internetwork. Simulations of small-scale magnetoconvection by, e.g., Stein \\& Nordlund (\\cite{stein_nordlund06}) show examples of granular-sized magnetic loops with similar behavior. They probably build up a sub-canopy of horizontal fields over granules (Wedemeyer-B\\\"{o}hm et al. \\cite{wedemeyerbohmetal09}). Surface dynamo simulations by Sch\\\"{u}ssler \\& V\\\"{o}gler (\\cite{schussler_vogler08}) and Steiner et al. (\\cite{steineretal08}) clearly show that this new component of the internetwork is a direct consequence of magnetic flux expulsion from the granular interior, not only to the intergranular lanes, but also to the upper photosphere. Cheung et al. (\\cite{cheungetal07}, \\cite{cheungetal08}) found in numerical simulations that emerging flux regions also contribute to the flux accumulation and its maintenance above granular upflows in the magnetic inversion layer. Our results show one particular example of this local process albeit terminated, probably due to magnetic reconnection. Better temporal and spatial resolution is needed for more detailed information on the process than our observations provide."
    },
    "0910/0910.3743_arXiv.txt": {
        "abstract": "{Spectroscopy of redshifted radio absorption of atomic and molecular species provide excellent probes of the cold component of the gas in the early Universe which can be used to address many important issues, such as measuring baryonic content, probing large-scale structure and galaxy evolution, as well as obtaining independent measurements of various combinations of fundamental constants at large look-back times. However, such systems are currently very rare with only 80 detected in \\HI\\ 21-cm and five in OH and millimetre-band species. Here we summarise the main selection criteria responsible for this and how the next generation of radio telescopes are expected to circumvent these through their wide instantaneous bandwidths and fields-of-view. Specifically: \\begin{enumerate} \\item {\\em \\HI\\ in absorbers occulting distant quasars:} These are usually found in known optical absorbers and wideband radio surveys could reveal a much fainter population. However, the 21-cm absorption strength may be correlated with the width of the singly ionised metal species, suggesting that these may be weak, and due to the effects of a flat expanding Universe on the covering factor, we expect the highest detection rates at $z < 1$.  \\item {\\em \\HI\\ absorption associated with the host galaxy/quasar:} Due to high degrees of ionisation/excitation rendering 21-cm undetectable near active nuclei with ultra-violet luminosities of $L_{\\rm UV}\\gapp10^{23}$ \\WpHz, future searches should be magnitude limited, e.g. at $z>1$, blue magnitudes should be $B \\gapp19$ with $z>2 \\Rightarrow B \\gapp21$ and $z>3 \\Rightarrow B \\gapp22$.  \\item {\\em OH (and millimetre-band) absorption:} For all of the known redshifted molecular absorption systems a correlation is found between the molecular fraction and the optical--near-infrared colour ($V-K$), with the five known OH absorbers all having $V-K\\gapp5$. Therefore spectral scans towards extremely red radio sources are expected to uncover any dusty intervening, molecular rich absorbers reponsible for the obscuration of the optical light.  \\end{enumerate} } \\FullConference{Panoramic Radio Astronomy: Wide-field 1-2 GHz research on galaxy evolution\\\\ June 2-5 2009\\\\ Groningen, the Netherlands}  \\begin{document} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.0786.txt": {
        "abstract": "A filament disappearance event was observed  on 22 May 2008  during our recent campaign JOP 178. The filament, situated in the southern hemisphere, showed sinistral chirality consistent with the hemispheric rule.  The event was well observed by several observatories in particular by THEMIS.   One day before the disappearance, H$\\alpha$ observations showed up and down flows in adjacent locations along the filament, which suggest plasma motions along twisted flux rope. THEMIS and GONG  observations show shearing photospheric motions leading to magnetic flux canceling around barbs. STEREO A, B spacecraft with separation angle 52.4 degrees, showed quite different views of this untwisting flux rope in He II 304 \\AA\\ images.  Here, we reconstruct the 3D geometry of the filament during its eruption phase using STEREO EUV He II 304 \\AA\\ images and find that the filament was highly inclined to the solar normal. The He II 304 \\AA\\ movies show individual threads, which oscillate and rise  to an altitude of about 120 Mm  with apparent velocities of about 100 km s$^{-1}$, during the rapid evolution phase. Finally, as the flux rope expands into the corona, the filament disappears by becoming optically thin to undetectable levels. No CME was detected by STEREO, only a faint CME was recorded by LASCO at the beginning of the disappearance phase at 02:00 UT, which could be due to partial filament eruption.  Further,  STEREO Fe XII 195 \\AA\\ images showed bright loops beneath the filament prior to the disappearance phase, suggesting  magnetic reconnection below the flux rope. ",
        "introduction": "\\label{S-Introduction} Two types of structures are commonly used to model filaments i.e., arcades structure or twisted flux tubes. Evidence of twisted flux tubes clearly appear during the eruptive phase of prominences \\cite{Gary2004}, \\cite{Torok2005}. Most quantitative models assume that the flux rope is in static equilibrium. 3-D magnetic models of filaments by extrapolating photospheric magnetograms into the corona have been developed \\cite{1998A&A...329.1125A}, \\cite{Aulanier2000}, \\cite{Dudik2008}, \\cite{aad2004}. Such models reproduce helical ropes overlying the polarity inversion line (PIL) and the filament plasma is assumed to be located in  dips of the helical field lines. In areas where magnetic parasitic polarity elements are located close to the PIL, the dips extend away from the main body of the filament creating barbs. It is consistent with the findings of the observers \\cite{Martin94},\\cite{Martin98} and \\cite{Wang2001}. One of the possible causes of filament eruption is instability %Instabilities in the filament equilibrium sometimes lead to  eruption \\cite{Forbes1991},\\cite{Isenberg1993}. Determination of true filament height is an important parameter in studies of filament eruption \\cite{Schrijver2008}.  Traditionally, this has been a difficult task and only via H$\\alpha$ observations of the limb prominence one could measure the filament height. The disadvantage of this method being: (i) it measures only the projected height, (ii) non-availability of magnetograms due to its location at the limb,  (iii) can not be used to follow height evolution for several days and (iv) no continuity of multi-temperature observations. Only recently, with the advent of stereoscopic observations by the twin spacecraft of the {\\it Solar Terrestrial Relations Observatory} (STEREO) mission called STEREO-A (Ahead) and STEREO-B (Behind), the true height of filament can be judged properly and the heating of the plasma can be tested \\cite{2008SSRv..136....5K};\\cite{2008SoPh..252..397G};\\cite{Liewer09}.  Here in this paper, we report on the multi-wavelength observations of a filament eruption event using ground as well as space based observatories during a joint observing campaign (JOP-178 from 20 to 25 May 2008). The filament was located in a large filament channel with very weak and diffuse polarities. The weak magnetic polarities in the filament channel are recognized using THEMIS/MTR instrument which has high polarimetric sensitivity and a simultaneous H$\\alpha$ scan. The evolution of these polarities is studied using full-disk GONG line-of-sight magnetograms which are available at a cadence of one minute. Further, we reconstruct the true filament height and eruption velocity using stereoscopic images by SECCHI/EUV instrument  in He II 304 \\AA\\ ($\\sim$60-80 $\\times 10^{3}$ K). The He II 304 \\AA\\ observations are very useful in tracing filaments because (i) the filament spine is much sharper and clearer \\cite{2007AAS...21012006M};\\cite{2007BASI...35..447J}, and (ii) the filament can be traced up to higher altitudes compared to H$\\alpha$ images \\cite{2007BASI...35..447J}.  With the help of stereoscopic observations by STEREO mission we reconstruct the filament geometry and  height during the phase.  It is difficult to derive the rise velocity by height-time profile as the filament becomes very diffuse during the disappearance phase, however, we try to estimate the projected rise velocity of individual threads from movies. Also, we identify possible reconnected loops below the flux rope. ",
        "conclusions": "A long  S-shaped filament composed of several segments and located in the Southern hemisphere was disappearing between May 20-22, 2008. This was observed during a coordinated campaign (JOP-178) involving ground-based instruments, THEMIS on the Canary Islands, MSDP at the Meudon solar tower, GONG magnetograph and H$\\alpha$ telescope in Udaipur, as well as space-based instruments SOHO/MDI and SECCHI/EUVI aboard STEREO. H$\\alpha$ instruments observed the progressive disappearance of the filament, segment after segment, between May 20 to May 22.  It was a long process with some impulsive phases. The last H$\\alpha$ segment was visible  on May 22 till 10:00 UT while the long spine of the S-shaped filament was observable in He II 304 \\AA\\ till 15:30 UT. Before disappearing filament segments showed high dynamics in H$\\alpha$ with blue and red shifts parallel to the filament axis.  MDI and GONG magnetograms show that the filament is located along the polarity inversion line between two weak magnetic field regions. The magnetic field in the filament channel was weak and rapidly changing. Local correlation tracking techniques applied to GONG polarities showed strong shear between these two regions two days before the eruption, then some vortex pattern and no noticeable motion after the eruption. THEMIS magnetograms allowed us to identify weak  minority polarities associated with the filament feet. The canceling flux/polarities in the vicinity of the feet were observed in the GONG movies a few hours before the disappearance of the associated segment.  Such cancelation of flux and disappearance of the foot-point barbs where the filament is tied to the photosphere could lead to  eruptions (Raadu {\\it et al.} 1987,1988).   The angular separation  between STEREO A and STEREO B spacecraft was 52.4 degrees during our observations, which gave quite different views of the filament. We used different methods to study the filament disappearance using He II 304 filtergrams: (i) study of filament dynamics and estimation of the projected rise velocity of thread-like structures using He II 304 movies, (ii)  by overlaying a spherical grid over the solar surface drawn over a sphere of 700 Mm to identify elevated features and filament feet and put an upper limit on the filament inclination, and (iii) using the triangulation algorithm SCC\\_MEASURE (part of STEREO data analysis library in Solarsoft package) to compute the altitude and inclination of the filament over the chromosphere during its disappearance phase. These  methods give consistent results.   The stereoscopic reconstructions using SECCHI/EUVI observations of 19 May 2007 filament eruption event were reported recently by \\cite{Liewer09} and \\cite{2008SoPh..252..397G}. While \\inlinecite{Liewer09} used scc\\_measure for reconstruction, \\inlinecite{2008SoPh..252..397G} used optical flow method to find displacements between features in stereoscopic pairs of SECCHI/EUVI 304 \\AA\\ images. The results of these two techniques are in good agreement \\cite{Liewer09}. These  results for the 19 May 2007 event showed that the filament eruption was  asymmetric and whip-like. The filament disappearance in our case is different from their study in two ways (i) filament was rooted in a weak diffuse bipolar magnetic region not in an active region, and (ii) there was no CME recorded for our filament disappearance event (however, one cannot rule out a CME as it could be below detection threshold). The length of the filament is too large spanning more than 10$^\\circ$ in latitude, as seen in He II 304 \\AA\\ images. We also use scc\\_measure to derive three-dimensional geometry of the filament which is used to derive its height and inclination. The filament was highly inclined to solar normal by about 47$^\\circ$. Such large inclinations are interesting as  theoretical models of filaments consider filament sheet as thin vertical slab.  The maximum filament height determined during the disappearance phase around 10:56 UT is estimated to be about 123 Mm in the middle portion of the filament. Determining height evolution of filament during the disappearance phase was difficult as it became quite diffuse, making identification of common features very difficult. The presence of fine threads running parallel to filament spine could be seen in movies. These threads are so tiny that it was difficult to track them over the chromosphere in the later phases, i.e., after around 11:00 UT. Only the EUVI movies allowed us to distinguish them over the chromosphere. The initial dense filament became an  untwisting flux rope with multiple threads with a fan-shaped structure which rose and disappeared one by one into the corona. The plasma becomes optically so thin as the flux rope rapidly expands in the corona that it remains no longer visible.  Bright structure observed in STEREO A and B images at 195 \\AA\\ is interpreted as reconnection loop system, located below the filament (flux-rope).  This bright structure is not visible earlier at 00:00 UT and appears during the onset of rapid disappearance phase at about 06:00 UT. The plasma of the filament is less and less dense as the flux rope  rises and expands. As the plasma is dispersed it is no more visible in filter 304 \\AA\\  nor in 195 \\AA\\.~ No CME was reported during this last phase (LASCO was not observing) by COR1 and COR2, the two coronagraphs aboard STEREO. The CME might be too faint to be detected or the magnetic field might be too weak to prevent the plasma's expansion to undetectable density levels.  Ground based data are very useful to understand the context of the event. Up to now we have the knowledge of eruption phenomena principally by using instruments with low temporal and/or spatial resolution. These observations are still largely used to study the association of filament eruption and CMEs. On the other hand high-cadence instrumentation with high-resolution have made a breakthrough in the nature of the dynamics of filament fine-structures. Non-eruptive filaments show counter-streaming flows along the fine structures and up-and-down flows in the filament-end (\\cite{Zirker98}, \\cite{Lin2005}, \\cite{Berger2008}). These velocities are lower than 10 km s$^{-1}$. With STEREO and ground-based observations we can find some relationship between the dynamics at large-scale and at small-scale. The H$\\alpha$ doppler velocities associated with the filament, that we measure in the present paper, are relatively small but they confirm previous observations of activated filaments with low spatial-resolution instruments (\\cite{Schmieder85a},\\cite{Schmieder85b}). The pair of aligned and elongated regions of oppositely directed velocities are interpreted in terms of a twisted  magnetic flux rope. The STEREO observations allow perspective effect to be removed in doppler velocities by taking into account the filament inclination. The corrected velocities are about 1.5 times more. The Doppler-shifts are also lower due to the computation by using the bisector method and not taking account the background chromosphere. In the paper by \\cite{Schmieder85a} and \\cite{Schmieder85b}, they could derive larger velocities by a factor 2 between the foot-points and of the order of 10 km s$^{-1}$ in the feet using a cloud model method. That is probably what we could expect with our present observations using the cloud model method. The standard method (bisector) indicates the general trend of the velocities with large uncertainty but indicates clearly that the filament is activated. The day before we did not observe such organized velocity pattern.  In the future, development of ground based instruments like dual-beam doppler imaging system, being developed at Udaipur Solar Observatory \\cite{Joshi2009} for detecting filament activation using high-cadence H$\\alpha$ dopplergrams,  and combined observations with STEREO will lead to further developments in our understanding of the filament eruptions.    \\begin{acks} The authors thank CNRS for allocating observing time on THEMIS. THEMIS is a French-Italian telescope installed at Observatorio El Teide, Tenerife, Spain. The observations from BBSO, GONG and SOHO are acknowledged. We thank V. Bommier for providing her UNNOFIT code for inversion of Stokes profiles. We also thank Arturo Lopez Ariste, G. Ruyman and  Cyril Delaigue for help in observations and data reductions. Further, SG acknowledges CEFIPRA funding for his visit to  Observatoire de Paris, Meudon, France under its project No. 3704-1. Also, SG acknowledges travel support to THEMIS, Tenerife by Observatoire de Paris, Meudon. We thank Fr$\\acute{\\rm e}$d$\\acute{\\rm e}$rique Auch$\\grave{\\rm e}$re for the STEREO movies. We also thank G. Molodij, J. Moity for the observations at the solar tower and P. Mein for helping in the MSDP data reduction. This work was supported by the European network SOLAIRE (MTRN\\_CT\\_2006\\_035484). \\end{acks}"
    },
    "0910/0910.5023_arXiv.txt": {
        "abstract": "Measurements of small-scale turbulent fluctuations in the solar wind find a non-zero right-handed magnetic helicity. This has been interpreted as evidence for ion cyclotron damping.  However, theoretical and empirical evidence suggests that the majority of the energy in solar wind turbulence resides in low frequency anisotropic kinetic \\Alfven wave fluctuations that are not subject to ion cyclotron damping.  We demonstrate that a dissipation range comprised of kinetic \\Alfven waves also produces a net right-handed fluctuating magnetic helicity signature consistent with observations. Thus, the observed magnetic helicity signature does not necessarily imply that ion cyclotron damping is energetically important in the solar wind. ",
        "introduction": "The identification of the physical mechanisms responsible for the dissipation of turbulence in the solar wind, and for the resulting heating of the solar wind plasma, remains an important and unsolved problem of heliospheric physics. An important clue to this problem is the observed non-zero fluctuating magnetic helicity signature at scales corresponding to the dissipation range of solar wind turbulence.  \\citet{Matthaeus:1982a} first proposed the ``fluctuating'' magnetic helicity as a diagnostic of solar wind turbulence, defining the ``reduced fluctuating'' magnetic helicity spectrum derivable from observational data (see \\S \\ref{sec:red} below).  A subsequent study, corresponding to scales within the inertial range, found values that fluctuated randomly in sign, and suggested an interpretation that ``a substantial degree of helicity or circular polarization exists throughout the wavenumber spectrum, but the sense of polarization or handedness alternates randomly'' \\citep{Matthaeus:1982b}. Based on a study of the fluctuating magnetic helicity of solutions to the linear Vlasov-Maxwell dispersion relation, \\citet{Gary:1986} suggested instead that, at inertial range scales, all eigenmodes have a very small {\\it intrinsic} normalized fluctuating magnetic helicity, eliminating the need to invoke an ensemble of waves with both left- and right-handed helicity to explain the observations.  Subsequent higher time resolution measurements, corresponding to scales in the dissipation range, exhibited a non-zero net reduced fluctuating magnetic helicity signature, with the sign apparently correlated with the direction of the magnetic sector \\citep{Goldstein:1994}. Assuming dominantly anti-sunward propagating waves, the study concluded that these fluctuations had right-handed helicity. The proposed interpretation was that left-hand polarized Alfv\\'en/ion cyclotron waves were preferentially damped by cyclotron resonance with the ions, leaving undamped right-hand polarized fast/whistler waves as the dominant wave mode in the dissipation range, producing the measured net reduced fluctuating magnetic helicity. We refer to this as the \\emph{cyclotron damping interpretation}.  A subsequent analysis of more solar wind intervals confirmed these findings for the dissipation range \\citep{Leamon:1998a}. \\citet{Leamon:1998b} argued that a comparison of the normalized cross-helicity in the inertial range (as a proxy for the dominant wave propagation direction in the dissipation range) to the measured normalized reduced fluctuating magnetic helicity provides evidence for the importance of ion cyclotron damping, which would selectively remove the left-hand polarized waves from the turbulence; using a simple rate balance calculation, they concluded that the ratio of damping by cyclotron resonant to non-cyclotron resonant dissipation mechanisms was of order unity. A recent study performing the same analysis on a much larger data set concurred with this conclusion \\citep{Hamilton:2008}.  In this Letter, we demonstrate that a dissipation range comprised of kinetic \\Alfven waves produces a reduced fluctuating magnetic helicity signature consistent with observations.  A dissipation range of this nature results from an anisotropic cascade to high perpendicular wavenumber with $k_\\perp \\gg k_\\parallel$; such a cascade is consistent with existing theories for low-frequency plasma turbulence \\citep{Goldreich:1995,Boldyrev:2006,Howes:2008a,Schekochihin:2009}, numerical simulations \\citep{Cho:2000,Howes:2008b}, and observations in the solar wind \\citep{Horbury:2008,Podesta:2009a}.  Our results imply that no conclusions can be drawn about the importance of ion cyclotron damping in the solar wind based on the observed magnetic helicity signature alone. ",
        "conclusions": "\\label{sec:discuss}  Predicting the values of $H_m^{'r}(\\omega')$ for solar wind turbulence based on equation (\\ref{eq:hmr}) requires understanding three issues: the scaling of the magnetic fluctuation spectrum with wavenumber, the imbalance of \\Alfven wave energy fluxes in opposite directions along the mean magnetic field, and the variation of the angle $\\theta$ between the solar wind velocity $\\V{v}$ and the mean magnetic field.  The 1-D magnetic energy spectrum in the solar wind typically scales as $k_1^{-5/3}$ in the inertial range and $k_1^p$ in the dissipation range, where $-2\\le p \\le -4$ \\citep{Smith:2006} and the effective wavenumber is $k_1=\\omega'/v$. It is clear from \\eqref{eq:hmr} that, when the plasma frame frequency $\\omega$ is negligible, the Doppler-shifted observed frequency always results in an effective wavenumber $k_1 \\le k$, with equality occurring only when the velocity $\\V{v}$ is aligned with the wave vector $\\V{k}$. We assume that, for homogeneous turbulence at the dissipation range scales, turbulent energy at fixed $k_\\perp$ and $k_\\parallel$ is uniformly spread over wave vectors with all possible angles $\\alpha$ about the mean magnetic field.  Because the fluctuation amplitude deceases for larger effective wavenumbers, the contribution to $H_m^{'r}(\\omega')$ is maximum at angle $\\alpha=0$; for angles $\\alpha$ yielding a Doppler shift to lower effective wavenumbers $k_1<(k_\\perp^2+k_\\parallel^2)^{1/2}$, the higher amplitude fluctuations at those lower wavenumbers will contribute more strongly to $H_m^{'r}(\\omega')$.  An accurate calculation of the magnetic helicity signature based on \\eqref{eq:hmr} must take into account the scaling of the magnetic energy spectrum.    \\begin{figure} \\resizebox{3.1in}{!}{\\includegraphics*[0.3in,2.in][8.0in,5.1in]{f2.eps}} \\caption{\\label{fig:mhelr} Normalized reduced fluctuating magnetic helicity $\\hat{\\sigma}^r_m(k_1)$ vs.~effective wavenumber $k_1$ due to a turbulent spectrum of kinetic \\Alfven waves with $\\theta=60^\\circ$. The solid line corresponds to the model 1-D energy spectrum while the dashed line corresponds to a $k^{-1}$ spectrum.} \\end{figure}  To compare to $\\sigma^r_m(k_1)$ derived from observations (for example, see Figure~1 of \\citet{Leamon:1998a}), we construct the normalized quantity \\begin{equation} \\hat{\\sigma}^r_m(k_1)= \\frac{\\sum_{\\V{k}} H'_m(\\V{k})\\frac{\\V{k}'\\cdot \\V{v} } {\\V{k}'\\cdot \\V{v}+\\omega} \\delta[\\omega'- (\\V{k}'\\cdot \\V{v}+\\omega)]} { \\sum_{\\V{k}} [|\\V{B}(\\V{k})|^2/k ]\\delta[\\omega'- (\\V{k}'\\cdot \\V{v}+\\omega)]}. \\label{eq:numsigm} \\end{equation} In evaluating \\eqref{eq:numsigm}, we assume a model 1-D energy spectrum\\footnote{On $150 \\times 150$ logarithmic gridpoints over $k_\\perp \\rho_i, k_\\parallel \\rho_i \\in [10^{-3},10^2]$, the model weights $B^2$ as a function of $k=({k_\\perp^2+k_\\parallel^2})^{1/2}$ using $B^2(k) = B_0^2 \\{[ (k \\rho_i)^{-1/3} + (k \\rho_i)^{4/3} ]/[1+(k \\rho_i)^{2}]\\}^2$.} that scales as $k^{-5/3}$ for $k\\rho_i \\ll 1$ and $ k^{-7/3}$ for $k\\rho_i \\gg 1$, consistent with theories for critically balanced turbulence \\citep{Goldreich:1995,Howes:2008b,Schekochihin:2009} and solar wind observations \\citep{Smith:2006}. In \\figref{fig:mhelr}, we plot $\\hat{\\sigma}^r_m(k_1)$ vs.~effective wavenumber $k_1=\\omega'/v$ for a turbulent spectrum filling the MHD \\Alfven and kinetic \\Alfven wave regimes ($k_\\perp>k_\\parallel$ and $k_\\parallel \\rho_i <1$) for $\\beta_i=1$, $T_i/T_e=1$, $v_{th_i}/c=10^{-4}$, $\\theta=60^\\circ$, and $v/v_A=10$. The contributions to $\\hat{\\sigma}^r_m(k_1)$ for all angles $\\alpha$ of each wave vector are collected in 120 logarithmically spaced bins in Doppler-shifted frequency.  The results are rather insensitive to the scaling of the 1-D magnetic energy spectrum over the range from $k^{-1}$ to $k^{-4}$. The solid line in \\figref{fig:mhelr} corresponds to the model spectrum assumed above, while the dashed line corresponds to a $k^{-1}$ energy spectrum. \\figref{fig:mhelr} demonstrates that turbulence consisting of \\Alfven and kinetic \\Alfven waves produces a positive (right-handed) magnetic helicity signature in the dissipation range at $k_1 \\rho_i \\gtrsim 1$.   The analysis presented in \\figref{fig:mhelr} considers only waves with $k_\\parallel >0$, so all of the waves in the summation in \\eqref{eq:hmr} are traveling in the same direction. If there were an equal \\Alfven wave energy flux in the opposite direction---a case of balanced energy fluxes, or zero cross helicity---the net $\\hat{\\sigma}^r_m(k_1)$ would be zero due to the odd symmetry of $H'_m(\\V{k})$ in $k_\\parallel$. It is often observed, at scales corresponding to the inertial range, that the energy flux in the anti-sunward direction dominates, leading to a large normalized cross helicity \\citep{Leamon:1998b}. If this imbalance of energy fluxes persists to the smaller scales associated with the dissipation range, a non-zero value of $\\hat{\\sigma}^r_m(k_1)$ is expected. However, theories of imbalanced MHD turbulence \\citep[][ and references therein]{Chandran:2008} predict that the turbulence is ``pinned'' to equal values of the oppositely directed energy fluxes at the dissipation scale. This implies that, at sufficiently high wavenumber $k_1$, the value of $\\hat{\\sigma}^r_m(k_1)$ should asymptote to zero. Thus, $\\hat{\\sigma}^r_m(k_1)$ in \\figref{fig:mhelr} would likely drop to zero more rapidly than shown, leaving a smaller positive net $\\hat{\\sigma}^r_m(k_1)$ around $k_1 \\rho_i \\sim 1$, consistent with observations \\citep{Goldstein:1994,Leamon:1998a,Hamilton:2008}.  We defer a detailed calculation of the effects of imbalance to a future paper.  The angle $\\theta$ between $\\V{B}_0$ and $\\V{v}$ is likely to vary during a measurement; this angle does not typically sample its full range $0 \\le \\theta \\le \\pi$, but has some distribution about the Parker spiral value. Calculations of $\\hat{\\sigma}^r_m(k_1)$ over $0 \\le \\theta \\le \\pi$ yield results that are qualitatively similar to \\figref{fig:mhelr}, so this averaging will not significantly change our results.  Taken together, we have demonstrated that a solar wind dissipation range comprised of kinetic \\Alfven waves produces a magnetic helicity signature consistent with observations, as presented in \\figref{fig:mhelr}.  The underlying assumption of the cyclotron damping interpretation of magnetic helicity measurements, an interpretation that dominates the solar wind literature \\citep{Goldstein:1994,Leamon:1998b,Leamon:1998a,Hamilton:2008}, is the slab model, $\\V{k}=k_\\parallel \\zhat$ and $k_\\perp=0$, i.e., purely parallel wave vectors.  As shown in \\figref{fig:mhel}, only in the limit $k_\\parallel \\gg k_\\perp$ does the \\Alfven wave root generate a left-handed helicity $\\sigma_m \\rightarrow -1$ as $k_\\parallel \\rho_i \\rightarrow \\sqrt{\\beta_i}$; in the same limit, the fast/whistler root generates a right-handed helicity $\\sigma_m \\rightarrow +1$ in a quantitatively similar manner (see Figure~9 of \\cite{Gary:1986}). Strong ion cyclotron damping of the Alfv\\'en/ion cyclotron waves as $k_\\parallel \\rho_i \\rightarrow 1$ \\citep{Gary:2004} would leave a remaining spectrum of right-handed fast/whistler waves, as proposed by cyclotron damping interpretation.  However, only if the majority of the turbulent fluctuations have $k_\\parallel \\gtrsim k_\\perp$ is the slab limit applicable, and only if significant energy resides in slab-like fluctuations are the conclusions drawn about the importance of cyclotron damping valid. There is, on the other hand, strong theoretical and empirical support for the hypothesis that the majority of the energy in solar wind turbulence has $k_\\perp \\gg k_\\parallel$ (see \\citealt{Howes:2008b} and references therein).  In this case, there is a transition to kinetic \\Alfven wave fluctuations at the scale of the ion Larmor radius.  This Letter demonstrates that a dissipation range comprised of kinetic \\Alfven waves produces a reduced fluctuating magnetic helicity signature consistent with observations."
    },
    "0910/0910.2617_arXiv.txt": {
        "abstract": "The {\\it Chandra} COSMOS Survey (C-COSMOS) is a large, 1.8 Ms, {\\it Chandra} program, that covers the central contiguous $\\sim$~0.92 deg$^2$ of the COSMOS field. C-COSMOS is the result of a complex tiling, with every position being observed in up to six overlapping pointings (four overlapping pointings in most of the central $\\sim$~0.45 deg$^2$ area with the best exposure, and two overlapping pointings in most of the surrounding area, covering an additional  $\\sim$~0.47 deg$^2$). Therefore, the full exploitation of the C-COSMOS data requires a dedicated and accurate analysis focused on three main issues: 1) maximizing the sensitivity when the PSF changes strongly among different observations of the same source (from $\\sim$~1 arcsec up to $\\sim~10$ arcsec half power radius); 2) resolving close pairs; and 3) obtaining the best source localization and count rate.  We present here our treatment of four key analysis items: source detection, localization, photometry, and survey sensitivity.  Our final procedure consists of a two step procedure: (1) a wavelet detection algorithm, to find source candidates, (2) a maximum likelihood Point Spread Function fitting algorithm to evaluate the source count rates and the probability that each source candidate is a fluctuation of the background. We discuss the main characteristics of this procedure, that was the result of detailed comparisons between different detection algorithms and photometry tools, calibrated with extensive and dedicated simulations. ",
        "introduction": "It is well known that X-ray surveys are an extremely efficient tool to select Active Galactic Nuclei (AGN). For example in the {\\it XMM-Newton} COSMOS survey, at the 0.5-2 keV limiting flux of 7$\\cdot$10$^{-16}$ erg s$^{-1}$ cm$^{-2}$, the AGN surface density is $\\sim$1000 deg$^{-2}$ (Hasinger et al. 2007, Cappelluti et al. 2007), a factor 2-4 greater than the AGN surface density in the most recent deep optical surveys, 250 deg$^{-2}$ in the COMBO-17 ( Wolf et al. 2003) and 470 deg$^{-2}$ in VVDS Survey (Gavignaud et al. 2006). There are four main causes for the higher efficiency of X-ray surveys in finding AGN: 1) X-rays directly trace the super massive black hole (SMBH) accretion, while AGN classification trough optical line spectroscopy may suffer of uncompleteness and/or misidentifications; 2) AGN are the dominant X-ray population. In fact most ($\\sim$ 80\\%) of the X-ray sources AGN in deep and shallow surveys turn out to be AGN, unlike at optical wavelengths. 3) 0.5-10 keV X-rays (the typical {\\it Chandra} and {\\it XMM-Newton} enery band) are capable to penetrate column densities up to $\\sim$10$^{24}$ cm$^{-2}$, allowing the selection of moderately obscured AGN; 4) low luminosity AGN are difficult to select in optical surveys, because their light is diluted in the host galaxy emission.  So far {\\it Chandra} and {\\it XMM-Newton} have performed several deep, pencil beam, and shallower but wider surveys. Fig. \\ref{surveys} compares the flux limit and area coverage of the main {\\it Chandra} and {\\it XMM-Newton} surveys. This figure shows that {\\it XMM-Newton} COSMOS and {\\it Chandra}-COSMOS (C-COSMOS, Elvis et al. 2009, Paper I hereafter) surveys are the deepest surveys on large contiguous area. The coverage of larger areas at similar flux limits is today achieved only by serendipitous surveys using mostly not contiguous areas (see e.g., CHAMP, Kim et al. 2004a, 2004b, Green et al. 2004).  \\begin{figure} \\begin{center} \\includegraphics[angle=0,height=8truecm]{xraysurveys.ps} \\caption{The 0.5-2 keV flux range vs. the area coverage for various surveys. The black solid lines represent few serendipitous surveys: Helllas2XMM (Baldi et al. 2002, symbol A), CHAMP (Kim et al. 2004a, 2004b, Green et al. 2004, symbol B), SEXSI (Harrison et al. 2003, symbol C), XMM-BSS (Della Ceca et al. 2004, symbol D), AXIS (Carrera et al. 2007, symbol E); the red dotted lines represent few deep pencil beam surveys: CDFN (Brandt et al. 2001, Alexander et al. 2003, symbol F), CDFS (Giacconi et al. 2001, Luo et al. 2008, symbol G), XMM-Newton Lockman Hole (Worsley et al. 2004, Brunner et al. 2008, symbol H); the blue dotted lines represent few wide shallow contiguous surveys: C-COSMOS (Elvis et al. 2009, symbol I), XMM-COSMOS (Hasinger et al. 2007, Cappelluti et al. 2007, 2009, symbol L), ELAIS-S1 (Puccetti et al. 2006, symbol M), E-CDF-S (Lehmer et al. 2005, symbol N), AEGIS-X (Laird et al. 2009, symbol O), SXDS (Ueda et al. 2008, symbol P). The black solid triangle represent the ROSAT all sky survey (RASS, Voges et al. 1999).} \\label{surveys} \\end{center} \\end{figure}  The Cosmic evolution survey (COSMOS, Scoville et al. 2007) is aimed at studying the interplay between the Large Scale Structure (LSS) in the Universe and the formation of galaxies, dark matter, and AGN. The COSMOS field is located near the equator (10h,+02degrees), covers $\\sim$~2 square degrees as originally defined by the HST/ACS imaging (Koekemoer et al. 2007), with subsequent deep and extended multi-wavelength coverage overlapping this area.  The size of COSMOS was chosen to sample LSS up to a linear size of about 50 Mpc h$^{-1}$ at z~$\\sim$~1-2, where AGN and star formation in galaxies are expected to peak. To study the role of AGN in galaxy evolution the X-ray data are fundamental. Therefore, the central square degree of the COSMOS field has been the target of a {\\it Chandra} ACIS-I, 1.8~Msec Very Large Program: the {\\it Chandra}-COSMOS survey.  The  C-COSMOS  survey  has  a  rather uniform  effective  exposure  of $\\sim~160$~ksec over a large area ($\\sim$~0.45 deg$^2$), thus reaching $\\sim$~3.5 times fainter fluxes than XMM-COSMOS in both 0.5-2 keV band and  2-7  keV band.  This  flux limit  is  below  the threshold  where starburst galaxies  become common in X-rays.  The  sharp {\\it Chandra} Point Spread  Function (PSF) allows  nearly unambiguous identification of optical counterparts (Civano et al. 2009, hereafter Paper III).{\\it Chandra} secures  the identifications of  X-ray sources down  to faint optical magnitude (i.e.,  I~$\\sim$~26), with only $\\sim$~2\\% ambiguous identifications,  significantly better  than the  $\\sim20\\%$ ambiguous identifications in XMM-Newton (Brusa et al. 2007).  The C-COSMOS survey has a complex tiling (see Fig. \\ref{overlay}) in comparison to other X-ray surveys, in which the overlapping areas of the single pointings are small and with similar PSFs (see e.g., the Extended Groth-Streep, AEGIS-X, Laird et al. 2009), or all the pointings are co-assial and nearly totally overlapping (see e.g., CDFS, Giacconi et al. 2001, Luo et al. 2008).  In the C-COSMOS tiling, the pointings are strongly overlapping and not-coassial. While this ensures a very uniform sensitivity over most of the field, each source is observed with up to six different PSFs, requiring the development of an analysis procedure for data observed with this mixture of PSF. The procedures presented in this paper are aimed at optimizing (1) source detection, (2) localization, (3) photometry, and (4) survey sensitivity. We have made detailed comparisons between different detection algorithms and photometry tools, testing them extensively on simulated data. We furthermore validate our results by detailed inspections of each single source candidate. Our final analysis consists of a two main steps:  \\begin{itemize}  \\item [1] \\underline{a wavelet detection algorithm}, {\\it PWDetect} (Damiani et al. 1997) is first used to find source candidates. This algorithm is optimized to cleanly separate nearby sources, to detect point-like sources on top of extended emission and to give the most accurate positions.  \\item [2] \\underline{A maximum likelihood PSF fitting algorithm} is then used to evaluate the source count rates and the probability that each source candidate is not a fluctuation of the background. We used the {\\it emldetect} algorithm (Cappelluti et al. 2007 and references therein). {\\it emldetect} works simultaneously with multiple overlapping pointings using PSFs appropriate to each one. This fitting method ensures accurate evaluation of the survey completeness and contamination, efficient deblending and good photometry for close pairs, which may be partly blended even at the {\\it Chandra} resolution. \\end{itemize}  As a third step, we also performed \\underline{aperture photometry} for each candidate X-ray source using 50\\%, 90\\%, and 95\\% encircled count fractions, using the PSFs appropriate to each observation. The aperture photometry is also used to check the results. Aperture photometry is preferable in all cases where the systematic error introduced by PSF fitting are larger than the statistical errors, i.e., for bright sources (count rates $\\geq 1$ counts/ksec).   The survey sensitivity is limited by both the net (i.e., including vignetting) exposure time, and by the actual PSF with which a given region of the area is observed. The latter issue is particularly relevant for the C-COSMOS tiling. We have developed an algorithm that evaluates the survey sensitivity at each position on C-COSMOS using a parameterization of the {\\it Chandra} ACIS-I PSF and taking into account the mixture of PSFs at each position. The resulting sensitivity maps have been compared and validated with extensive simulations.  The paper is organized as following: in Sect. 2 we briefly present the C-COSMOS observations and data reduction; we describe the simulations in Sect. 3; how they were used to select the most efficient detection algorithm and the final source characterization procedure is described in Sect. 4; the completeness and reliability are shown in Sect. 5; in Sect. 6 we apply this procedure to the observed data; in Sect. 7 we present the calculation of survey sensitivity, the sky-coverage, and X-ray number counts using the simulated data. Finally, in Sect. 8 we compare C-COSMOS to a similar {\\it Chandra} survey, i.e., AEGIS-X, and in Sect. 9 we give our conclusion. ",
        "conclusions": "The complex tiling of C-COSMOS survey required the development of a tailored multistep procedure to fully exploit the data. Detailed simulations were used to test different detection (sliding cell and wavelet) and photometry (PSF fitting and aperture photometry) algorithms. In particular, we compared the results obtained using the SAS {\\it eboxdetect} and {\\it emldetect} tasks, used for the XMM-COSMOS survey (Cappelluti et al. 2007, 2009), with those obtained using the {\\it PWDetect} code (Damiani et al. 1997).  Through these tests we selected a procedure consisting in first identifying source candidates using {\\it PWDetect}, and then performing accurate PSF fitting photometry and evaluating aperture photometry for each source candidate.  In this way we obtained subarcsec source localizations and accurate photometry even for partly blended sources.  We set a threshold for source detection to $P=2\\cdot 10^{-5}$, which implies a completeness of 87.5\\% and 68\\% for sources with at least 12 and 7 F band counts, respectively, and 3 to 5 spurious detections in the F band at the same count limits, respectively.  We evaluated the survey sensitivity and the sky-coverage, through an analytical method, tuned using simulations.  We then evaluated the $\\log N$ -- $\\log S$ of the detected sources in the simulations down to F, S, and H band flux limits of F$_{x}$$\\sim~2.3\\cdot 10^{-16}$, $\\sim~1.6\\cdot 10^{-15}$, and $\\sim~9.6\\cdot 10^{-16}$ erg s$^{-1}$ cm$^{-2}$, respectively.  Finally we compared the C-COSMOS survey to the AEGIS-X survey, a {\\it Chandra} survey with similar sky-coverage and total exposure time, but using non overlapping ACIS-I pointings.  We found that the complex tiling of C-COSMOS helps in obtaining a contiguous area with uniform sensitivity and somewhat higher source density. The overlap of several pointings with different PSF at the same position produces an effective source extraction region of $\\sim$~3 arcsec radius, i.e., significantly wider than the Chandra PSF at off-axis angles smaller than 5-6 arcmin. This produces a number of independent detection cells per unit area smaller than in a single ACIS-I pointing survey like AEGIS-X, which in turn implies a smaller number of spurious sources at each given detection threshold."
    },
    "0910/0910.1614_arXiv.txt": {
        "abstract": "An analysis of archival mid-infrared (mid-IR) spectra of Seyfert galaxies from the \\textit{Spitzer Space Telescope} observations is presented. We characterize the nature of the mid-IR active nuclear continuum by subtracting a template starburst spectrum from the Seyfert spectra. The long wavelength part of the spectrum contains a strong contribution from the starburst-heated cool dust; this is used to effectively separate starburst-dominated Seyferts from those dominated by the active nuclear continuum. Within the latter category, the strength of the active nuclear continuum drops rapidly beyond $\\sim 20\\mum$. On average, type 2 Seyferts have weaker short-wavelength active nuclear continua as compared to type 1 Seyferts. Type 2 Seyferts can be divided into two types, those with strong poly-cyclic aromatic hydrocarbon (PAH) bands and those without. The latter type show polarized broad emission lines in their optical spectra. The PAH-dominated type 2 Seyferts and Seyfert 1.8/1.9s show very similar mid-IR spectra. However, after the subtraction of the starburst component, there is a striking similarity in the active nuclear continuum of all Seyfert optical types. PAH-dominated Seyfert 2s and Seyfert 1.8/1.9s tend to show weak active nuclear continua in general. A few type 2 Seyferts with weak/absent PAH bands show a bump in the spectrum between 15 and 20$\\mum$. We suggest that this bump is the peak of a warm ($\\sim$200$\\kel$) blackbody dust emission, which becomes clearly visible when the short-wavelength continuum is weaker. This warm blackbody emission is also observed in other Seyfert optical sub-types, suggesting a common origin in these active galactic nuclei. ",
        "introduction": "The complete mid-infrared (mid-IR) spectra of Seyfert galaxies \\citep[a class of active galactic nuclei (AGN):][]{1943ApJ....97...28S} have been available only in recent past \\citep[\\eg][]{2002A&A...393..821S,2005SSRv..119..355V,2005ApJ...633..706W,2007ApJ...671..124D}. Many {\\it Spitzer} spectra of nearby Seyfert galaxies show a strong contribution from star-forming features in the form of poly-cyclic aromatic hydrocarbon (PAH) bands \\citep[\\eg][]{2000A&A...357..839C,2006AJ....132..401B,2008ApJ...676..836T}. A number of these features, such as the 7.7$\\mum$ and the 17$\\mum$ PAH complex, are strongly blended with each other, complicating the estimation of the underlying active nuclear continuum. Further, the mid-IR opacity includes a strong contribution from the silicate bands at $10$ and $18\\mum$. Thus, estimating the intrinsic active nuclear continua in the mid-IR is a non-trivial task, with the primary hurdle being the subtraction of the starburst component.  Previous studies have hinted that the continuum in the 1--8$\\mum$ range is non-stellar in origin and is likely a result of thermal emission from dust heated close to sublimation temperature by the optical/ultra-violet continuum from the central source \\citep{1986ApJ...308...59E,2001AJ....121.1369A,2004ApJ...614..122I,2008arXiv0807.4695M}. In a multi-wavelength photometric study of spectral energy distributions (SED) of predominantly radio-quiet Sloan Digital Sky Survey (SDSS) quasars, \\citet{2006ApJS..166..470R} and \\citet{2007ApJ...661...30G} noted that the 1--8$\\mum$ spectral index ($\\alpha_{\\nu}$) is strongly anti-correlated with infrared luminosity in type 1 quasars. The more luminous quasars have flatter 1--8$\\mum$ slopes. A linear correlation between the optical continuum luminosity and the infrared luminosity suggests that the observed bump around $\\sim2.2\\mum$ in the SED is driven by the dust re-emission. For example, the $2.2\\mum$ bump is clearly visible in the near-IR spectrum of Mrk~1239 \\citep{2006MNRAS.367L..57R}.  Due to the presence of strong PAH bands in the 5--8$\\mum$ range in many Seyfert spectra, direct measurement of the active nuclear continuum in this region is only possible in sources with very weak or absent PAH features. An alternative is to subtract the starburst contribution using a template starburst spectrum and then study the residual continuum. We take this simple approach in this paper. In Section~2, we describe our archival sample and the data analysis techniques. We discuss the observed mid-IR spectra in Section~3. In Section~4, we use simple continuum diagnostics that allows us to classify Seyfert mid-IR spectra into PAH-dominated and AGN-dominated groups and understand continuum properties. In Section~5, we discuss the starburst contribution in Seyfert spectra and the resulting continuum shapes after subtraction of the starburst template spectrum. In Section~6, we summarize our results. ",
        "conclusions": "An analysis of archival {\\it Spitzer Space Telescope} mid-IR spectra of Seyfert galaxies is presented. We focus on understanding the intrinsic shape of the active nuclear continuum in the mid-IR region and how it relates to other properties of the source such as the 10$\\mum$ silicate optical depth. We assumed a template spectrum for the starburst component, and subtracted it from the Seyfert spectra to separate the active nuclear contribution from the circum-nuclear starburst contribution. Our primary conclusions from this study are as follows: \\begin{enumerate} \\item Seyfert spectra are classified effectively between AGN- and starburst-dominated categories based on the spectral indices, $\\alpha_{\\lambda}(5.5\\textrm{--}14.7\\mum)$ and $\\alpha_{\\lambda}(20\\textrm{--}30\\mum)$ (see Figure~\\ref{fig:continua}). Seyferts dominated by the AGN contribution have flatter spectra with $\\alpha_{\\lambda}(5.5\\textrm{--}14.7\\mum) \\sim -1.13$. The added starburst contribution from the host galaxy in the large {\\it Spitzer} aperture (see Table~\\ref{tab:objdata}) makes $\\alpha_{\\lambda}(20.0\\textrm{--}30.0\\mum)$ more positive or steeper. A key property that distinguishes AGN-dominated spectra is that the 20--30$\\mum$ continuum is flatter ($\\alpha_{\\lambda}\\sim -2$) than the very steep 20--30$\\mum$ continuum of starburst-dominated objects ($\\alpha_{\\lambda}\\sim 1$). The 20--30$\\mum$ continuum is likely formed from the Rayleigh--Jeans tail of the ``warm'' dust component. It is likely that there are multiple ``warm'' components with different temperatures. Further, Type 2 Seyferts with polarized broad emission lines in their optical spectra (type S1h) show $\\alpha_{\\lambda}(5.5\\textrm{--}14.7\\mum) \\sim - 0.49$ much different than average type 1 Seyferts that show $\\alpha_{\\lambda}(5.5\\textrm{--}14.7\\mum) \\sim -1.13$. Note the steeper and weak short-wavelength continuum in Figure~\\ref{fig:all_spec_nosb}, as compared to type 1 Seyferts. This is a direct evidence for presence of the dust torus that blocks our view of the hot dust closer in. \\item After starburst subtraction, Seyfert 1.8/1.9s and Seyfert 2s with strong PAH features in their spectra show similar active nuclear continuum as Seyfert 2s with weak/absent PAH features in their mid-IR spectra and polarized broad emission lines in their optical spectra (see Figure~\\ref{fig:all_spec_nosb}). This suggests presence of similar quantities and/or properties of dusty material around the central accretion disk in these type 2 sources. \\citet{2003ApJ...583..632T} had proposed existence of two types of Seyfert 2s: the HBLR and the non-HBLRs. We compared spectral indices in the 5--8$\\mum$ region after starburst subtraction and find that both the HBLR and non-HBLR show similar spectral indices (Figure~\\ref{fig:after-sb-sub}, bottom), suggesting similarity rather than differences between the two classifications. While, non-HBLRs tend to show stronger starburst contribution as compared to their AGN contribution, the converse (that starburst-dominated systems lack BLR signatures) is not necessarily true. The additional host galaxy contribution likely complicates the identification of BLR in these systems. As we show in Figure~\\ref{fig:continua}, simple continuum indices effectively separate AGN-dominated Seyferts from starburst-dominated Seyferts in the mid-IR. \\item \\citet{2007ApJ...671..124D} showed that a large part of the silicate absorption in some Seyfert galaxies comes from starburst-heated cold dust in the host galaxy rather than the dust torus. This conclusion was based on the fact that only highly inclined galaxies showed large silicate optical depths. Here, we put that result on a better statistical basis by presenting a correlation between the $b/a$ and the measured optical depth for this sample of 109 sources (see Figure~\\ref{fig:opt-depth-ba}). All objects with significant optical depth are highly inclined and are also classified as Seyfert 2s or Seyfert 1.8/1.9s. This confirms previous results by \\cite{1980AJ.....85..198K} and \\cite{1995ApJ...454...95M}, and highlights the importance of considering the host galaxy contribution in concealing AGN. \\item On average, Seyfert galaxies dominated by the AGN continuum tend to show weak silicate absorption ($\\tau_{9.7} \\lesssim 0.4$). The short wavelength continuum index ($\\alpha_{\\lambda}(5.5\\textrm{--}14.7\\mum)$) and the apparent silicate optical depth $\\tau_{9.7}$ (Figure~\\ref{fig:opt-depth}) may to be correlated in AGN-dominated objects. Seyfert optical types form a continuous sequence of increasing optical depth along this correlation from type 1s, type 1.8/1.9s, to type 2s with HBLRs. This validates the general inclination dependence inherent in AGN models noted before in the mid-IR by \\citet{2007ApJ...655L..77H}. But, as can be noted in Figure~\\ref{fig:opt-depth}, there is not a strict relationship between the strength of silicate features and the optical spectral type. \\item Figure~\\ref{fig:after-sb-sub} (top) shows that there are at least two types of dust distributions in the active nuclear region, one that generates the short-wavelength continua and another that generates the long-wavelength continua. The major difference between these two dust distributions is likely to be their mean temperatures. By subtracting the starburst and associated cool dust contribution, we have essentially removed the starburst component that contributes most to the variety in Seyfert spectra \\citep{2006AJ....132..401B}. In Figure~\\ref{fig:bump}, we demonstrate the existence of this ``warm'' component which dominates the long-wavelength continuum, by separating it from the hot dust component in Mrk 766. This simple template subtraction exercise provides proof that Seyfert spectra are primarily thermal in nature and composed of at least three thermal components with $T\\sim$ 1000, 200, and 60$\\kel$. The reported ``break'' in the spectrum at $\\sim$ 20$\\mum$ in type 1 Seyfert spectra is a result of this warm component being brighter at $\\sim 20 \\mum$ than the Rayleigh--Jeans tail of the hot component. The above-mentioned similarity of continua beyond $\\sim15\\mum$ in AGN-dominated sources is also due to this warm dust component being present in almost all observed AGN spectra. \\end{enumerate}"
    },
    "0910/0910.3611_arXiv.txt": {
        "abstract": "In a dynamical-radiative model we recently developed to describe the physics of compact, GHz-Peaked-Spectrum (GPS) sources, the relativistic jets propagate across the inner, kpc-sized region of the host galaxy, while the electron population of the expanding lobes evolves and emits synchrotron and inverse-Compton (IC) radiation. Interstellar-medium gas clouds engulfed by the expanding lobes, and photoionized by the active nucleus, are responsible for the radio spectral turnover through free-free absorption (FFA) of the synchrotron photons. The model provides a description of the evolution of the spectral energy distribution (SED) of GPS sources with their expansion, predicting significant and complex high-energy emission, from the X-ray to the $\\gamma$-ray frequency domain. Here, we test this model with the broad-band SEDs of a sample of eleven X-ray emitting GPS galaxies with Compact-Symmetric-Object (CSO) morphology, and show that: (i)  the shape of the radio continuum at frequencies lower than the spectral turnover is indeed well accounted for by the FFA mechanism; (ii) the observed X-ray spectra can be interpreted as non-thermal radiation produced via IC scattering of the local radiation fields off the lobe particles, providing a viable alternative to the thermal, accretion-disk dominated scenario. We also show that the relation between the hydrogen column densities derived from the X-ray (\\NHmath) and radio (\\NHImath) data of the sources is suggestive of a positive correlation, which, if confirmed by future observations, would provide further support to our scenario of high-energy emitting lobes. ",
        "introduction": "\\label{sec_introduction}  The power released by active galactic nuclei (AGNs) is currently interpreted in terms of conversion of gravitational energy to radiative energy by accretion processes feeding the central supermassive black hole (BH) with environmental gas. The triggering, maintenance, and fading of the AGN activity, as well as the link of these processes with the physical conditions of the environment, from sub- to super-galactic scales, are still widely debated issues, and keep on stimulating a large variety of scientific investigations. In this context, radio galaxies are ideal laboratories, because they offer an edge-on view of both their nuclei and their relativistic jets, launched from the galactic center and reaching up to Mpc distances. The variety of powers, sizes, and morphologies displayed by radio galaxies not only provides pieces of evidence of different environmental physical conditions, but also samples subsequent stages of the source evolution. In particular, key sources for the investigation of the very first phases of the evolution of radio galaxies are the Gigahertz-Peaked-Spectrum (GPS) sources associated with galaxies and characterized by a Compact Symmetric Object (CSO) radio morphology.  GPS sources \\citep[see][for a review]{odea1998} are a class of powerful radio sources ($P_{1.4\\, {\\rm GHz}}\\gtsim 10^{25}$ \\wattperhz) displaying convex radio spectra that turn over at frequencies of about 0.5--1 GHz; they  make up a conspicuous fraction ($\\sim$10\\%) of the radio sources found in high-frequency radio surveys, and are optically identified with either galaxies or quasars; their radio sizes are smaller than about 1 kpc, and their radio morphologies reveal either core-jet structures or mini-lobes embedding terminal hotspots and possibly straddling a central core. CSOs \\citep{wilkinson1994,readhead1996} have instead emerged in VLBI surveys as a class of radio sources with compact ($\\ltsim$500 pc) and symmetric radio structures, making them resemble ``classical double'' radio galaxies in miniature. Whereas there is no debate on the true compactness of CSOs, the compactness of a fraction of the core-jet GPS sources might be the result of foreshortening of extended radio sources roughly aligned with the line of sight; especially (although not exclusively) in these cases, the GPS itself might be a transient spectral state: flux-density variability and polarization studies at different radio frequencies are thus instrumental to the selection of bona-fide samples of GPS sources \\citep{tinti2005,torniainen2005,torniainen2007,orienti2008a}. The overlap between the GPS-source and the CSO classes is however significant: GPS sources associated with galaxies are most likely to display a CSO morphology, and most CSOs exhibit a GPS. This overlap is explained in terms of synchrotron radio spectra dominated by the emission of the mini-lobes, and suffering from absorption effects causing the turnover about 1 GHz \\citep{snellen2000}. GPS/CSOs {associated with} galaxies are thus high-confidence candidates for truly compact sources.  Although compact GPS sources were proposed to be young objects soon after their discovery \\citep[e.g.,][]{shklovsky1965}, a widely discussed alternative to the youth scenario was the frustration scenario, ascribing the small source size to confinement by a particularly dense interstellar medium (ISM) preventing the jet expansion \\citep{vanbreugel1984}. Because the required ISM confining densities are not confirmed by recent studies \\citep[e.g.,][]{morganti2008}, the confinement scenario is considered less likely, although it might still apply to selected objects \\citep[e.g.,][]{garciaburillo2007}. Much observational evidence has instead accumulated in favour of the youth scenario, the most compelling measurement being the detection, in several GPS/CSO galaxies, of hotspot advance velocities about 0.1--0.2$c$ \\citep{owsianikconway1998,owsianik1998,tschager2000,taylor2000,gugliucci2005}, which indicate kinematical source ages not higher than a few $10^3$ years, in good agreement with spectral ageing estimates \\citep{murgia1999}. Further evidence of youth comes from the underluminosity of GPS/CSO optical narrow-emission lines, suggesting that the Str\\\"omgren sphere of the recently-triggered AGN is still in an expansion phase in these sources, and we are thus witnessing the birth of their Narrow-Line Region \\citep[NLR;][]{vink2006}.  The youth scenario is part of a wider evolutionary scenario \\citep{fanti1995,snellen2000}, according to which GPS/CSOs would first evolve into symmetric Compact-Steep-Spectrum (CSS) sources, equally powerful but with sizes of about 1--15 kpc, and then further expand outside the host galaxy to become large-scale, powerful radio galaxies. The evolutionary link between GPS and CSS sources is supported by their membership to the anticorrelation between linear size and peak frequency \\citep[$LS \\propto \\nu_{\\rm p}^{-0.65}$;][]{odea1997}. However, it is not clear yet whether all GPS/CSOs will eventually become large-scale radio sources, and whether the evolved sources will display a Fanaroff-Riley type I (FRI) or type II (FRII) morphology, because of the frequency's observational limits in the exploration of this relation for super-galactic sized sources, as well as because of the biases in the surveys from which the distribution of the radio sources in the power-linear size ($P-LS$) diagram is drawn \\citep[e.g.,][and references therein]{snellen2000}.   The first systematic studies of the host galaxies of GPS sources, conducted in the optical and near-infrared (NIR) bands \\citep{snellen1996,odea1996,devries1998a,devries1998b,devries2000} showed that they are, as CSS sources, characterized by luminous masses, brightness profiles, and optical-NIR colors more typical of passively- or non-evolving giant elliptical galaxies, than either spirals, small ellipticals, brightest cluster galaxies (BCGs), or central-dominant (cD) galaxies at similar redshifts. This seems to be consistent with what was found for a sample of FRIIs rather than with the properties of FRIs, which are preferentially hosted by cD galaxies \\citep{owen1989,owen1991,zirbel1996}, and suggest that GPS sources are more likely to evolve, through the CSS phase, in FRIIs rather than in FRIs \\citep{devries2003}. However, the  morphologies of the majority \\citep[$\\sim$60\\%;][]{odea1996} of the hosts of GPS sources show signs of recent mergers and/or interactions. This evidence, apparently at odds with the passively-evolving scenario, was first interpreted as an indication that GPS sources might be associated with the first of a sequence of mergers \\citep{devries2003}, possibly leading to the formation of a BCG or a cD only over time scales much longer than the lifetime of the radio source. In subsequent investigations, however, the presence of young stellar population in these giant ellipticals was revealed by means of stellar-population synthesis models applied to photometric \\citep{devries2007} and spectrophotometric data \\citep{holt2009}. Furthermore, \\citet{devries2007} showed the similarity, out to $z \\simeq 0.55$, between the luminosities of the GPS-source hosts and of the Luminous Red Galaxies (LRGs), which are thought to represent the most massive early-type galaxies, and are often associated with BCGs. Therefore, the properties of the hosts do not provide unambiguous indications on the subsequent GPS/CSO evolutionary stages.  Despite these results, population studies highlight the existence of far too many compact sources compared to the number of powerful large-scale objects \\citep{odea1997}. Beside the possibility, not easy to justify, that these sources undergo a short-lived activity \\citep{readhead1996}, two scenarios prove to be particularly attractive to explain this statistical evidence. In the first scenario, the jet activity is intermittent on time scales comparable to the age of the small sources \\citep[e.g.,][]{reynolds1997}: this scenario, supported by the observations of double-double radio galaxies \\citep{lara1999,schoenmakers2000,kaiser2000}, and by the detection of candidates for dying compact sources \\citep{giroletti2005,parma2007}, finds a natural explanation in the framework of accretion-disk instabilities \\citep{czerny2009}. An accretion disk operating above a given threshold accretion rate ($\\dot m_* \\simeq 0.025\\, \\dot m_{Edd}$) is expected to experience instabilities driven by the radiation pressure, causing the source undergo alternate phases of high and low activity; for moderate accretion rates, the duration of the active phases is short enough ($\\sim 10^3-10^4$ years) to make the compact source unable to grow to super-galactic sizes, whereas accretion rates close to the Eddington limit are required for the development of large-scale sources \\citep{czerny2009}. In the second scenario, many sources undergo processes of jet-flow turbulent disruption before their lobes can grow to large sizes, either disappearing from the radio sky \\citep{alexander2000} or developing an FRI-type morphology \\citep{kaiser2007}, in both cases experiencing a drop in luminosity. The recent finding that FRIs smaller than $\\sim$40 kpc were almost absent in a sample of radio galaxies optically identified in the SDSS \\citep{best2009} seems to suggest that all radio sources begin their life with collimated jets, but less powerful jets are more easily disrupted, as the source grows, in denser environments, leading to the production of FRI sources. Further evidence supporting this view might derive from the confirmation that the large fraction of low-power, compact objects with non-FRI morphology discovered within a sample of FRI candidates at $1<z<2$ \\citep{chiaberge2009}, are indeed young sources. The two above scenarios are not necessarily in conflict with each other: for sufficiently high accretion rates, accretion-disk instabilities might partake to the radio-source evolution, otherwise driven by the interplay between jet power and environmental conditions. However, this picture needs to be further explored.   It has recently become clear that instrumental clues on the evolution of young GPS sources may come from the X-ray electromagnetic window. Early detections of GPS sources by {\\it ASCA} \\citep{odea2000} triggered extensive searches for X-ray emission from these sources; thanks to the capabilities of {\\it XMM-Newton} and {\\it Chandra}, several GPS/CSOs are now known to be strong X-ray emitters \\citep{guainazzi2004,vink2006,guainazzi2006,siemiginowska2008,tengstrand2009}. However, the origin of this X-ray emission is not known yet.  The best spatial resolution currently available in the X-ray band ($\\sim$1$''$ with {\\it Chandra}) is not sufficient to resolve the X-ray morphology of most GPS/CSOs: thus, the identification of the origin of the X-ray emission in these objects relies entirely on spectral studies. The components of GPS/CSOs first proposed to be the source of the observed X-rays are the ISM of the host galaxy shocked by the expanding radio lobes \\citep{heinz1998,odea2000}, and the accretion disk's hot corona \\citep{guainazzi2004,vink2006,guainazzi2006,siemiginowska2008}: both components are expected to generate {\\it thermal} X-ray radiation, but they are difficult to disentangle and do not rule out alternative scenarios \\citep[][and references therein]{siemiginowska2009}. Furthermore, the location of GPS/CSOs on the radio/X-ray luminosity plane does not fully elucidate the relationship between ``young'' and ``mature'' radio galaxies: GPS/CSOs are as X-ray luminous as FRIIs, despite their larger radio luminosity, but they seem to lie on the high radio-power tail of the radio-core/X-ray correlation discovered for FRIs \\citep[][and references therein]{tengstrand2009}.  A way of investigating the origin of the X-rays in GPS/CSOs, and address, at the same time, some of the open issues on the physics and fate of young radio sources, is to simultaneously model the dynamical and radiative evolution of GPS/CSO galaxies, and constrain the models through the wealth of multiwavelength data currently available. We recently proposed a dynamical-radiative model that, for the first time, describes the evolution of the broad-band emission of a young GPS with a given jet kinetic power, as it expands through the ISM of the host galaxy \\citep{stawarz2008}. The model accounts for the radiative contribution of the various AGN components, as well as for environmental absorption effects, and predicts significant and complex {\\it non-thermal} X-ray to $\\gamma$-ray emission.  In the present paper, we apply the model to a sample of GPS/CSO galaxies, with the manifold aim of reproducing their complex radio to X-ray spectra, and constraining relevant source parameters as the jet kinetic power, the accretion rate, and the absorption mechanisms. New observational evidence supporting the model is also discussed. The paper is organized as follows: in Section \\ref{sec_model}, we briefly review the main features of the dynamical-radiative model we presented in \\citet{stawarz2008}; in Section \\ref{sec_sample}, we describe the source sample that we chose for the application of the model, and the relevant broad-band data; in Section \\ref{sec_SEDmodeling}, we show the results of the SED modeling; in Section \\ref{sec_support}, we present further evidence supporting the proposed scenario; we discuss our results in Section \\ref{sec_discussion}, and in Section \\ref{sec_conclusions} we draw our conclusions.  A model of $\\Lambda$-dominated universe \\citep[$\\Omega_\\mathrm{{\\Lambda}}=0.7$, $\\Omega_{\\mathrm{M}} = 0.3$, $\\Omega_{\\mathrm{k}} = 0$;][]{spergel2003}, with $H_{0} = 72$ km~\\persec~Mpc$^{-1}$ \\citep{freedman2001} will be adopted throughout this paper. Source-intrinsic quantities from the literature given below were corrected for this cosmology, unless otherwise stated. ",
        "conclusions": "\\label{sec_conclusions} For the first time, we systematically compared the broad-band spectra of a sample of GPS/CSO galaxies with a model. Specifically, we showed that the dynamical-radiative model we proposed in \\citet{stawarz2008} can reproduce the observed emission, from the radio to the X-ray energy range, of eleven  GPS/CSO galaxies characterized by different linear sizes, and thus sampling different stages of the source expansion.  The radio spectra were modeled as synchrotron emission by a lobe electron population that represents the evolution of an injected hot-spot electron population suffering from adiabatic and radiative energy losses. At frequencies below the turnover, the shape of most of the spectra could be satisfactorily accounted for by an absorption scenario dominated by the FFA mechanism.  The X-ray emission could be interpreted as non-thermal radiation produced through IC scattering of the local thermal radiation fields off the lobe electron population. Among the possible seed photon fields, the dominant role in the generation of the X-ray radiation via IC scattering is played by the MFIR radiation, which we associated to the putative circumnuclear dusty torus; dust associated with circumnuclear star-forming regions might also contribute to this emission, but this would not affect our results significantly. The UV radiation emitted by the accretion disk gives a contribution comparable to that of the MFIR photons in a few sources only, but provides the bulk of the seed photons responsible for the $\\gamma$-ray emission. Our model proved to be a viable alternative to the thermal, accretion-disk dominated scenario for the interpretation of the X-ray emission of GPS/CSO galaxies.  Finally, from a more observational perspective, our radiative model finds further support in the comparison of the measurements of the radio and X-ray hydrogen column densities of a sub-sample of our source ensemble: the data suggest a positive correlation, which, if confirmed, would point towards the co-spatiality of the radio and X-ray emission regions. New radio measurements, necessary to improve the statistics of our correlation sample, as well as the quality of the available data, are currently being performed."
    },
    "0910/0910.5421_arXiv.txt": {
        "abstract": "{ The majority of the observed planetary nebulae exhibit elliptical or bipolar structures. Recent observations have shown that asymmetries already start during the last stages of the AGB phase. Theoretical modeling has indicated that magnetically collimated jets may be responsible for the formation of the non-spherical planetary nebulae. Direct measurement of the magnetic field of evolved stars is possible using polarization observations of different maser species occurring in the circumstellar envelopes around these stars.}{The aim of this project is to measure the Zeeman splitting caused by the magnetic field in the OH and H$_2$O maser regions occurring in the circumstellar envelope and bipolar outflow of the evolved star W43A. We compare the magnetic field obtained in the OH maser region with the one measured in the H$_2$O maser jet.}{We used the UK Multi-Element Radio Linked Interferometer Network (MERLIN) to observe the polarization of the OH masers in the circumstellar envelope of W43A. Likewise, we used the Green Bank Telescope (GBT) observations to measure the magnetic field strength obtained previously in the H$_2$O maser jet.} {We report a measured magnetic field of approximately 100 $\\mu G$ in the OH maser region of the circumstellar envelope around W43A. The GBT observations reveal a magnetic field strength $B_{\\rm ||}$ of $\\sim$30 mG changing sign across the H$_2$O masers at the tip of the red-shifted lobe of the bipolar outflow. We also find that the OH maser shell shows no sign of non-spherical expansion and that it probably has an expansion velocity that is typical for the shells of regular OH/IR stars.} {The GBT observations confirm that the magnetic field collimates the H$_{2}$O maser jet, while the OH maser observations show that a strong large scale magnetic field is present in the envelope surrounding the W43A central star. The magnetic field in the OH maser envelope is consistent with the one extrapolated from the H$_2$O measurements, confirming that magnetic fields play an important role in the entire circumstellar environment of W43A.} ",
        "introduction": "At the end of their evolution, low mass stars undergo a period of high mass loss ($\\dot{M} \\propto 10^{-4} -  10^{-7}$  M$_{\\odot}$/yr) that is important for enriching the interstellar medium with processed molecules. During this stage the star climbs up the asymptotic giant branch (AGB)  in the Hertzsprung-Russell (H-R) diagram. AGB stars are generally observed to be spherically symmetric \\citep{habingreview}. However, planetary nebulae (PNe), supposedly formed out of the ejected outer envelopes of AGB stars, often show large departures from spherical symmetry. The origin and development of these asymmetries is not clearly understood. Observations of collimated jets and outflows of material in a number of PNe have been reported \\citep[e.g][]{sahai1998, miranda2001, alcolea2001}.  \\cite{sahai1998} propose that the precession of these jets is responsible for the observed asymmetries. These jets are likely formed when the star leaves the AGB and undergoes a transition to become a PN \\citep[e.g.][]{sahai1998,imainature,miranda2001}.  Magnetic fields can play an important role in shaping the circumstellar envelope (CSE) of evolved stars and can produce asymmetries during the transition from a spherical symmetric star into a non-spherical PN. They are also possible agents for collimating the jets around these sources \\citep{garcia2005}. It is not clear how stars can maintain a significant magnetic field throughout the giant phases, as the drag of a large scale magnetic field would brake any stellar dynamo if no additional source of angular momentum is present \\citep[e.g.][]{soker1998}. However, theoretical models have shown that AGB stars can generate the magnetic field through a dynamo interaction between the fast rotating core and the slow rotating envelope \\citep{blackmannature}. Alternatively, the presence of a heavy planet or a binary companion as the additional source of angular momentum can maintain the magnetic field \\citep[e.g.][]{frank2004}.  Polarization observations of different maser species in the CSE of these stars provide a unique tool for understanding the role of the magnetic field in the process of jet collimation. Strong magnetic fields have been observed throughout the entire CSE of these stars for different molecular species (e.g. \\citealt{etoka2004, wouternature,herpin2006} using OH, H$_{2}$O and SiO masers, respectively).  \\citet{likkel1992} introduced a separate class of post-AGB sources where H$_{2}$O maser velocity range exceeds that of OH maser in the CSE. High resolution H$_{2}$O maser observations trace highly collimated jets in the inner envelopes of these objects. These objects, the so-called water fountain sources, are thought to be in the transition stage to PNe. An archetype of this class is W43A, located at a distance of 2.6 kpc from the sun \\citep{diamond1985}, for which the H$_2$O masers have been shown to occur at the tips of a strongly collimated and precessing bipolar jet \\citep{imainature}. Polarization observations of the H$_2$O masers of W43A reveal a strong magnetic field apparently collimating the jet \\citep{wouternature}.   Here we describe observations of the Zeeman splitting of OH masers in the CSE and H$_{2}$O masers in the jet of W43A. Our aim is to investigate the magnetic field strength and morphology in the low density region of the CSE of this object, which is in transition to become a PN and confirm the magnetic field in the H$_2$O maser jet. The observations are described in \\S \\ref{observations} and the results of the  observed spectra and the measured Zeeman splitting are presented in \\S \\ref{results}. The results are discussed in \\S \\ref{discussion}, where we investigate the observed Zeeman splitting and the significance of other non-Zeeman effects. This is followed by conclusions in \\S \\ref{conclusion}. ",
        "conclusions": "The non-spherical shape of PNe is thought to be related to outflows already generated during the AGB phase. Magnetic fields are considered as collimating agents of the jets around evolved stars. The magnetic field and jet characteristics of W43A, an evolved star in transition to a PN, have been previously reported from H$_{2}$O maser polarization observations. Our GBT observations reveal a magnetic field strength $B_{\\rm ||}$ changing sign from $31\\pm8$~mG to $-24\\pm7$~mG across the H$_2$O masers in the red-shifted lobe of the W43A precessing jet. We observed the OH  maser region of the CSE of this star and measured circular and linear polarization. Due to Faraday rotation, we can not determine the magnetic field configuration in the OH maser shell. However, the measured circular polarization, which is attributed to the Zeeman effect, implies a magnetic field of 100 $\\mu$G in the OH  maser region of W43A. Our result is consistent with the predicted magnetic field extrapolated from the blue-shifted H$_{2}$O  maser region of W43A, and further confirms that the magnetic field plays an important role in the transition from a spherical AGB star to  a non-spherical PN.    We considered the motion and expansion of the OH maser region of W43A. The observed expansion is  2.67$\\pm$1.37 mas/yr, which corresponds to an expansion velocity of $\\sim$18 km/s. While the significantly offset red- and blue-shifted caps of the OH maser shell indicates it is  a-spherical, the measured OH maser motions do not show any signs of a clear bipolar expansion in the plane of the sky."
    },
    "0910/0910.0032_arXiv.txt": {
        "abstract": "A number of synchronous moons are thought to harbor water oceans beneath their outer ice shells.  A subsurface ocean frictionally decouples the shell from the interior. This has led to proposals that a weak tidal or atmospheric torque might cause the shell to rotate differentially with respect to the synchronously rotating interior.  Applications along these lines have been made to Europa and Titan.  However, the shell is coupled to the ocean by an elastic torque. As a result of centrifugal and tidal forces, the ocean would assume an ellipsoidal shape with its long axis aligned toward the parent planet.  Any displacement of the shell away from its equilibrium position would induce strains thereby increasing its elastic energy and giving rise to an elastic restoring torque. In the investigation reported on here, the elastic torque is compared with the tidal torque acting on Europa and the atmospheric torque acting on Titan.  Regarding Europa, it is shown that the tidal torque is far too weak to produce stresses that could fracture the ice shell, thus refuting an idea that has been widely advocated. Instead, it is suggested that the cracks arise from time-dependent stresses due to non-hydrostatic gravity anomalies from tidally driven, episodic convection in the satellite's interior.  Two years of Cassini RADAR observations of Titan's surface have been interpreted as implying an angular displacement of $\\sim0.24$ degrees relative to synchronous rotation. Compatibility of the amplitude and phase of the observed non-synchronous rotation with estimates of the atmospheric torque requires that Titan's shell be decoupled from its interior. We find that the elastic torque balances the seasonal atmospheric torque at an angular displacement $\\lesssim0.05$ degrees, effectively coupling the shell to the interior. Moreover, if Titan's surface were spinning faster than synchronous, the tidal torque tending to restore synchronous rotation would almost certainly be larger than the atmospheric torque. There must either be a problem with the interpretation of the radar observations, or with our basic understanding of Titan's atmosphere and/or interior. ",
        "introduction": "Synchronous, or near synchronous, spinning satellites with ice shells and subsurface oceans are an intriguing class of solar system bodies. Part of their appeal is the possibility that the oceans could support life. Currently Jupiter's satellite Europa and Saturn's satellite Titan are the prime candidates for membership in this class.  The induced magnetic field of Europa as measured by the magnetometer on the Galileo spacecraft essentially proves the presence of a current-carrying, near-surface liquid layer, i.e. a salty subsurface ocean \\citep{2000Sci...289.1340K}.  All mechanisms proposed for creating the observed morphologies of cracks on Europa's surface require the presence of a near-surface fluid layer  \\citep{1998Icar..135...25G,1999Icar..141..263G,1999Sci...285.1899H}. Most mechanisms also invoke an immeasurably slow, super-synchronous spin of the ice shell due to tidal torques associated with Europa's orbital eccentricity. The eccentricity, currently 0.01, is continually excited by resonant interactions with Io and Ganymede \\citep{Peale_etal79, OjakangasStevenson86}. \\cite{1984Icar...58..186G} suggest that tidal torques could drive a super-synchronous rotation of Europa. \\cite{1998Natur.391..368G} argue that the presence of a subsurface ocean might allow the surface ice shell to rotate non-synchronously even if the core is locked in synchronous rotation.  Models of Titan's thermal evolution are consistent with the presence of a subsurface ocean. Moreover, motions of Titan's surface features as tracked by the Cassini spacecraft's radar have been interpreted as implying a slow, non- synchronous rotation \\citep{Stiles_etal08,Stiles_etal09,Stiles_etal10}.   \\cite{2008Sci...319.1649L} propose that the non-synchronous rotation is driven by a seasonally changing atmospheric torque acting on an ice shell decoupled from the interior by a subsurface ocean.  The plan of our paper is as follows. As a result of centrifugal and tidal forces, the interiors of synchronous satellites have ellipsoidal figures whose longest axes point toward the parent planet.  We derive these forces and resultant equilibrium shape in \\S  \\ref{sec:potentialtheory}.  In \\S  \\ref{sec:energies}, we evaluate the increase in elastic and gravitational energy associated with the angular displacement of an elastic outer shell away from its equilibrium position and derive the elastic torque.  In \\S \\ref{sec:tidaltorque}, we outline a derivation of the tidal torque exerted on a synchronous shell due to the satellite's orbital eccentricity.  We derive the angular displacements at which the elastic torque balances the tidal torque acting on Europa's shell and the atmospheric torque acting on Titan's shell in \\S \\ref{sec:applications}.  The results of this exercise imply that the long cracks on Europa do not result from non-synchronous rotation of the ice shell driven by the tidal torque. They also cast doubt on claims of non-synchronous rotation for Titan.  We devote \\S \\ref{sec:gravityanomalies} to an explanation for the formation of large cracks on Europa's surface due to geoid variations associated with episodic mantle convection driven by tidal heating.  We end with a short discussion in \\S\\ref{sec:discussion}.  The main points of our investigation can be appreciated with order of magnitude estimates.  Consequently more detailed considerations are relegated to a series of Appendices. ",
        "conclusions": "\\label{sec:discussion}  We conclude with a few parting thoughts.  Our calculation of the elastic torque follows from the assumption that the shell's deviatoric stress vanishes in its tidally deformed state.   But there is another possibility.\\footnote{As pointed out to us by both Re'em Sari and Jay Melosh.}   The shell might have formed rotating super-synchronously with a deviatoric stress that would vanish if it were to assume an axisymmetric shape. Then non-synchronous rotation could proceed without any variation of elastic energy and thus unimpeded by an elastic torque. Although this is a logical possibility, the weakness of the drivers of non-synchronous rotation for Europa and Titan render it implausible.  The tiny values of $\\lambda_t$ calculated in \\S \\ref{sec:applications} are a measure of just how axisymmetric the unstressed shells would have to be in order to sustain non-synchronous rotation.  We have demonstrated that Europa's long cracks are not opened by stresses due to non-synchronous rotation driven by the torque arising from the diurnal tide.  Although this removes much of the {\\it evidence} for non-synchronous rotation, it does not preclude it.  Provided the small stresses estimated in \\S \\ref{sec:applications} could relax on timescale $t_{\\rm relax}$, nonsynchronous rotation {\\it might} proceed on timescale $t_{\\rm relax}/\\lambda_t$. Our use of might instead of would is deliberate. If the shell were to spin faster than synchronous, it would be subject to a component of tidal torque arising from the semi-diurnal tide which would tend to restore synchronous spin.  In the limit of small orbital eccentricity, the semi-diurnal component of tidal torque would overwhelm the diurnal component unless the tidal $Q$ were a steeply increasing function of period.   Non-hydrostatic moment differences would drive the ice shell to reorient with respect to a system of axes pointing toward the planet, along the orbit, and normal to it. However, because the ellipsoidal shape of the ocean surface would maintain its orientation in these axis, the elastic torque would oppose this type of polar wander \\citep{Willemann84}. Significant reorientation of the shell would require that the ice either crack or creep.  Provided Titan's shell is $\\sim100$km thick, the presence of an elastic torque presents a challenge to the interpretation of Titan's non-synchronous rotation.  Observational evidence for the non-synchronous rotation of Titan cannot be lightly dismissed. Although we suspect an error in the data analysis, it is also possible that we are missing something important in our understanding of the properties of Titan's atmosphere and/or interior.  \\appendix*"
    },
    "0910/0910.0204_arXiv.txt": {
        "abstract": "{} {With the aim of increasing the sample of M31 clusters for which a colour-magnitude diagram is available, we searched the HST archive for ACS images containing objects included in the Revised Bologna Catalogue of M31 globular clusters \\thanks{RBC Version 3.5 available at: www.bo.astro.it/M31}.} {Sixty-three such objects were found. We used the ACS images to confirm or revise their classification and were able to obtain useful CMDs for 11 old globular clusters and 6 luminous young clusters. We obtained simultaneous estimates of the distance, reddening, and metallicity of old clusters by comparing their observed field-decontaminated CMDs with a grid of template clusters of the Milky Way. We estimated the age of the young clusters by fitting with theoretical isochrones.} {For the old clusters, we found metallicities in the range $-0.4\\le $[Fe/H]$\\le -1.9$. The individual estimates generally agree with existing spectroscopic estimates. At least four of them display a clear blue horizontal branch, indicating ages $\\ga 10$ Gyr. All six candidate young clusters are found to have ages $<1$ Gyr. The photometry of the clusters is made publicly available through a dedicated web page.} {With the present work the total number of M31 GCs with reliable optical CMD increases from 35 to 44 for the old clusters, and from 7 to 11 for the young ones. The old clusters show similar characteristics to those of the MW. We discuss the case of the cluster B407, with a metallicity [Fe/H]$\\simeq -0.6$ and located at a large projected distance from the centre of M31 (R$_p=19.8$ kpc) and from the major axis of the galaxy (Y$=11.3$ kpc). Metal-rich globulars at large galactocentric distances are rare both in M31 and in the Milky Way. B407, in addition, has a velocity in stark contrast with the rotation pattern shared by the bulk of M31 clusters of similar metallicity. This, along with other empirical evidence, supports the hypothesis that the cluster (together with B403) is physically associated with a substructure in the halo of M31 that has been interpreted as the relic of a merging event. } ",
        "introduction": "%   Over the past $\\sim$20 years, the globular cluster (GC)  system of M31 has been the subject of intensive study   both from the ground and from space-borne observatories (see Rich et al. 2005; Galleti et al. 2004 - hereafter G04, 2006a, 2007; Huxor et al. 2008;  Lee et al. 2008  and  Caldwell et al. 2009 - hereafter C09, for recent reviews and references). One of the main aims of these studies was to collect as much as possible information on  the GCs in M31 and compare it with our knowledge of the GCs in the Galaxy,  so as to derive better insight into the formation and (chemical and dynamical)  evolution of these two spiral galaxies and possibly of galaxies in general. The advent of the Hubble Space Telescope provided the unprecedented opportunity to obtain colour-magnitude  diagrams (CMD) of M31 clusters, thus adding a completely new perspective to this research.  Substantial contributions in this field have been made by many  investigators. At present, sufficiently accurate visual CMDs for a meaningful comparison  with their Galactic counterparts have been published for 35 GCs in M31. Except for one that was observed from the ground (MGC1, Martin et al. 2006), a good fraction of these have been obtained with the HST-WFPC2 (Ajhar  et  al. 1996,  Rich et al. 1996, Fusi Pecci et al. 1996, Holland et al. 1997, Jablonka et al.  2000, Meylan et al. 2001, Rich  et al. 2005)  until the better resolution and sensitivity of the ACS allowed even  more accurate CMDs at fainter limiting magnitudes (Brown et al. 2004;  Huxor et al. 2004, 2005, 2008; Galleti et al. 2006b; Mackey et al. 2006, 2007).  In addition to photometric  quality, which is essential for the analysis  of individual objects, a good statistical coverage is also important for a better understanding of the GC system.  To increase the sample of M31 GCs with a CMD of individual member stars,  we searched the HST archive for ACS images of objects that are listed in the Revised Bologna Catalogue of M31 clusters (RBC, see G04). We found useful ACS images containing 69 such objects (see Fig.~\\ref{f:m31}). The retrieved material allowed us to confirm or revise the classification of all of them and to obtain CMD of individual stars for 17, 11 likely old globulars, and 6 young luminous clusters (like those discussed in Williams \\& Hodge 2001; Fusi Pecci et al 2005 and Perina et al. 2009a). This paper is devoted to the analysis of these data.  In Sect. 2 we present the target list, and in Sect. 3 we describe the adopted reduction procedures that yielded the CMDs.   Section 4 is devoted to describe the method we have used to  estimate the  metallicity, reddening, and distance from each individual CMD for which a sufficiently reliable decontamination from the non-member field components  was feasible. In Sect. 5 specific notes and comments on the results are presented for each  of the 11 GCs (the primary targets) and for the other objects for which a sufficiently meaningful photometry was carried out. In Sect. 6 we discuss a possible connection between a few clusters and a large substructure recently found in M31. Finally, Sect. 7 contains some general considerations and conclusions. ",
        "conclusions": "%  We have analysed 63 objects listed in the RBC  for which HST/ACS images were publicly available in the HST Archive. We have confirmed or revised their classification based on the inspection of these images and we were able to obtain meaningful CMD for 11 likely old GCs and 6 young bright clusters.  We estimated distance, reddening, and metallicity for the eleven old GCs, by comparing the field-decontaminated CMD of the clusters with a grid of ridge lines of well-studied template clusters of the Milky Way. Our reddening and metallicity estimates are, in general, in satisfactory agreement with previous independent measures.  As reported in Table \\ref{t:parme}, we have also determined for each cluster an estimate of the magnitude level of the HB measured on the HB ridge line of the  reference GGC that best fits the observed CMD, with a typical error of $\\pm$0.15 mag.  One of the clusters of our sample (B407) is identified as a possible member of a large sub-structure recently found in the halo of M31, and interpreted as a relic of past (minor) merging episodes (NE Shelf; Ibata et al. 2001b, 2005, 2007; Richardson et al. 2008). The cluster has a metallicity that is much higher than the bulk of M31 clusters residing at the same (large) distance from the M31 centre/major axis, and its kinematics is very different from the large majority of M31 GCs having similar metallicity. The GC B403 (not included in our sample) is found to share the same properties and is also indicated as a possible member of the NE Shelf.   We estimated the age also for six candidate young clusters,  by comparing their observed CMD with theoretical isochrones.  We confirm that all of them are younger than 1 Gyr, in good agreement with previous studies.   With the present analysis the total number of M31 confirmed GCs with published reliable optical CMDs increases from 35 to 44 for the old globulars, and from 7 to 11 for the young bright ones (BLCCs).    The photometric catalogues of the clusters studied here will be made publicly available through a dedicated web page\\footnote{\\tt http://www.bo.astro.it/M31/hstcatalog}."
    },
    "0910/0910.5567_arXiv.txt": {
        "abstract": "%  A detailed chemical composition analysis based on a high-resolution ($R \\simeq 35,000$) CCD spectrum is presented for a newly discovered post-AGB star in the globular cluster M79 for the first time. The elemental abundance results of M79 Post-AGB star are found to be $[C/Fe]\\simeq$-0.7, $[O/Fe]=$+1.4, $[\\alpha - process/Fe]\\simeq$0.5, and $[s-process/Fe]\\simeq$-0.1. The surprising result is that the iron abundance of the star is apparently about 0.6 dex less than that of the cluster's red giants as reported by published studies including a recent high-resolution spectroscopic analysis by Carretta and colleagues. ",
        "introduction": "%  \\noindent In this study, published recently in full by \\c{S}ahin \\& Lambert (2009), we report on an abundance analysis of the A-type m79 PAGB star  discovered by Siegel \\& Bond (2009, in preparation) in the globular cluster M79 and compare its composition to that of the cluster's red giants. The initial mass of this star must have been slightly in excess of the mass of stars now at the main sequence turn-off, say, $M \\simeq 0.8M_\\odot$. The star's composition may be referenced to that of the cluster's red giant stars for which abundance analyses have been reported. Comparison of abundances for the PAGB and RGB stars may reveal changes imposed by the evolution beyond the RGB; such changes are not necessarily attributable exclusively  to internal nucleosynthesis and dredge-up. It was in the spirit of comparing the compositions of the PAGB and RGB stars that we undertook our analysis. For the RGB stars, we use results kindly provided in advance of publication by Carretta (2008, private communication). \\vskip 0.2 cm   \\begin{figure} \\centering \\includegraphics[width = 20cm,,height=115mm, angle=90]{PAPER_spectrum_part_NEW.ps} \\caption{The spectrum for the PAGB star over the wavelength regions between 4467-4500 \\AA\\ (upper panel) and 4500-4530 \\AA\\ (lower panel). Selected lines are identified.} \\end{figure} ",
        "conclusions": ""
    },
    "0910/0910.3341_arXiv.txt": {
        "abstract": "We calculate durations and spectral parameters for 207 Swift bursts detected by the BAT instrument from April 2007 to August 2009, including 67 events with measured redshifts.  This is the first supplement to our catalog of 425 Swift GRBs (147 with redshifts) starting from GRB~041220. This complete and extensive data set, analyzed with a unified methodology, allows us to conduct an accurate census of intrinsic GRB energetics, hardnesses, durations, and redshifts. The GRB world model we derive reproduces well the observables from both Swift and pre-Swift satellites. Comparing to the cosmic star formation rate, we estimate that only about 0.1\\% of massive stars explode as bright GRBs. There is strong evidence for evolution in the Swift population at intermediate and high-z, and we can rule out (at the 5-sigma level) that this is due to evolution in the luminosity function of GRBs. Instead, the Swift sample suggests a modest propensity for low-metallicity, evidenced by an increase in the rate density with redshift. Treating the multivariate data and selection effects rigorously, we find a real, intrinsic correlation between $E_{\\rm iso}$ and $E_{\\rm pk}$ (and possibly also $T_{r45,z}$); however, the correlation {\\it is not} a narrow log-log relation and its observed appearance is strongly detector-dependent. We also estimate the high-z rate ($3-9$\\% of GRBs at z beyond 5) and discuss the extent of a large missing population of low-$E_{\\rm pk,obs}$ XRFs as well as a potentially large missing population of short-duration GRBs that will be probed by EXIST. ",
        "introduction": "\\label{sec:intro}  The {\\it Swift}~satellite \\citep{gehrels04} has transformed the study of Gamma-ray Bursts (GRBs) and their afterglows.  Our knowledge of the early X-ray afterglows has increased tremendously due to the dramatic success of the X-ray Telescope \\citep[XRT;][]{burrows05}. However, our understanding of the prompt emission properties has lagged.  This is due in part to the narrow energy bandpass of the Burst Alert Telescope \\citep[BAT;][]{bart05}, which precludes direct measurement of the broad GRB spectra and tends to weaken any inferences about the $\\nu F_{\\nu}$ spectral peak energy $E_{\\rm pk,obs}$ and the bolometric GRB fluence.  In the first installment of our ``Complete BAT Catalog of Swift GRBs and Spectra'' \\citep[][hereafter Paper I]{butleretal07}, we treat these limitations of the BAT in a statistically rigorous fashion and study tantalizing pre-Swift correlations between the host-frame characteristics of GRBs \\citep[e.g.,][]{lpm00,fr00,nmb00,shaf03,amati02,lamb04,ggl04,firmani06}. A number of these potential log-log relations appear dramatically different in the Swift-era sample, with a broader scatter, and a shift in normalization toward the detector threshold. From this, we concluded that the origin of these correlations was tied more closely to the detection process than to the intrinsic physics of GRBs.  We present here fits to the lightcurves and spectra of additional Swift GRBs detected between April 17, 2007 and August 13, 2009, nearly doubling the overall sample.  As summarized in Paper I and below, our analysis is extremely uniform (nearly fully-automated), with well-defined survey flux limits,  and our fits allow for detailed propagation of errors.  These features are critical to the estimate of GRB rates, the focus of the current work.  Our primary goal is to uncover, with realistic estimates of uncertainty, the intrinsic GRB production rate as a function of redshift $\\dot \\rho(z)$.  To measure this quantity, it is necessary to model the GRB luminosity function $\\phi(L)$ allowing for possible intrinsic and detection-based correlations.  We derive a model that describes both Swift and pre-Swift rates well as a function of hardness, duration, flux, and redshift (Section \\ref{sec:redux}). As we discuss in Section \\ref{sec:fitting} below, we find a highly significant, intrinsic correlation between $E_{\\rm iso}$ and $E_{\\rm pk}$; however, the observed correlation has large scatter and is strongly instrument-dependent. Figure \\ref{fig:logN-logS} shows how the overall number of events observed by Swift --- and not just the correlation between observed quantities --- requires us to build temporal and spectral dependences into modelling the luminosity function.  In the next section, we summarize the utility and historical background behind making a plot like Figure \\ref{fig:logN-logS}. We then discuss in Section \\ref{sec:fitting} the models that successfully recreate the curves in Figure \\ref{fig:logN-logS} and also the implications for the GRB luminosity function (Section \\ref{sec:energetics}) and comoving rate density (Section \\ref{sec:rho}). In Section \\ref{sec:exist} we make self-consistent predictions for the expected observed redshift distribution for all Swift GRBs --- including the 60\\% fraction of Swift GRBs for which a spectroscopic $z$ has not been measured --- and also for the planned and more-sensitive EXIST experiment. We expect EXIST to thrive with respect to the detection of both short and long duration GRBs at high redshift.  \\begin{figure} \\hspace{-0.2in} \\includegraphics[width=3.7in]{logN-logS.pdf} \\caption{\\small Strong spectral and temporal dependence in the number of GRBs with effective count rate $C_{\\rm eff}$ (Section \\ref{sec:overview}) above a given value (i.e., the Swift ``logN--logS'' curve; see Section \\ref{sec:prior}).  We plot here long duration ($T_{90}>3$ s) GRBs with hardness above and below the median Swift $E_{\\rm pk,obs}= 100$ keV and also short duration ($T_{90}<3$ s) GRBs.  The rate of long-duration hard GRBs is turning over at low flux levels, while the rate of long-duration soft GRBs rises more strongly.  This is a gradual effect in $E_{\\rm pk,obs}$. Although the logN--logS slope for long-duration GRBs does not appear to be duration dependent, the Swift short-duration GRB population is strongly rising in number to low flux levels, showing no significant sign of a turn-over. The curves expected without a cutoff --- derived in Section \\ref{sec:fitting} --- due to the detector are plotted as dotted lines. The dashed red curve (barely visible) at the left of the short-duration curve accounts for the detector threshold following the non-parametric prescription of \\citet{pet96}. } \\label{fig:logN-logS} \\end{figure}  \\subsection{Prior Rate and Luminosity Function Estimates} \\label{sec:prior}  There is a rich literature describing optimal ways of counting GRBs to determine their distance and intrinsic flux. In the pre-afterglow era, counting focused on the observed flux distributions. The number of events $N$ with observed flux greater than $S$ --- the so called ``logN--logS'' curve --- showed early evidence (over many decades in $S$) for slope $S^{-3/2}$ expected for a homogeneous, isotropic, and static source population in a Euclidean universe \\citep[HISE; e.g.,][]{hurley91,hs90}.   A powerful statistic for examining the source counts is $V/V_{\\rm max} = (C/C_{\\rm min})^{-3/2}$ \\citep{schmidt68}, a measure of the volume probed by a source detected with $C$ counts relative to a possible minimum number of observable counts $C_{\\rm min}$.  The expectation is that $\\langle V/V_{\\rm max} \\rangle = 0.5$ for HISE.  The BATSE experiment provided the first strong evidence from a single experiment \\citep{meegan92} --- a deficit of low $S$ GRBS and $\\langle V/V_{\\rm max} \\rangle <0.5$ --- for a departure from homogeneity, while the spatial counts showed clearly an isotropic population. To study whether these modest departures imply that GRBs are very local (a Galactic halo population) or cosmological required examination beyond the first moment $\\langle V/V_{\\rm max} \\rangle$ in the $V$ (or $C$) distribution  \\citep[e.g.,][]{band92,ht93,pet93}.  The first GRB redshifts \\citep[e.g.,][]{metzger97} defined a cosmological origin and a vast energy release. Connecting the small number of GRBs with $z$ to the large population of GRBs without $z$ required, in general, careful modelling of and strong assumptions for the intrinsic luminosity and number density distributions in order to reproduce the observed flux data \\citep[e.g.,][]{piran92, piran99, cp95, fb95, lw95, lw98, hh97, schmidt99, schmidt01, sb01, guetta05}. Exceptions to the parametric approach were studies utilizing luminosity criteria (i.e., possible correlations of observables with luminosity) to derive ``pseudo-redshifts'' for the full GRB sample \\citep[e.g.,][]{norris00,fr00,schaf01,lloyd02,mur03,yon04,firm04,kocevski06,schmidt09}.  These studies generally found a rising GRB rate to $z \\lessim 2$, similar to the cosmic star formation rate \\citep[e.g.,][]{madau96}, but potentially continuing to remain flat or even rising to $z\\sim 12$.  Notably, some of these works \\citep[e.g.,][]{lloyd02,yon04}, using hazard statistics \\citep{lb71,ep92,pet93,mp99}, also found evidence for potential strong luminosity evolution, parameterized as $L\\propto (1+z)^a$, with $a$ in the range 1.5--2.5. The luminosity function itself appears to generally be characterized well as broken powerlaw, with a break at $L\\sim 10^{51-52}$ erg s$^{-1}$ and a flat or slowly rising slope to low-energies, strongly dependent upon the instrumental detection model.  The connection between GRBs and the deaths of massive stars (now firmly established, e.g., \\citet{stanek03,hjorth03}; see \\citet{woosbloom06} for a review) sped progress by motivating an assumption that the GRB rate follows star formation \\citep[e.g.,][]{wijers98,lr00,pm01,cs02,bloom03,gor04,nat05}. Very recently, thanks to Swift and the impressive efforts of ground-based observers, a growing sample of GRBs with spectroscopic redshifts has allowed for direct tabulation of GRB intrinsic luminosities \\citep[e.g.,][]{kocevskibutler08} and redshifts \\citep[e.g.,][]{jak05,jak06}.  The large number of redshifts has also enabled a detailed comparison of the intrinsic GRB rate to the cosmic star formation rate \\citep[SFR, e.g.,][]{daigne06,salv07,kistler08,kistler09,salv09a,salv09b}.  Perhaps the most intriguing, shared feature of these studies is a strong indication of evolution in the GRB population.  Above $z\\approx 2$ and possibly extending to $z\\approx 8$, the Swit GRB rate is increasing far faster than star formation \\citep[e.g.,][]{kistler08,kistler09}, and it is not clear to what extent this is due to GRBs in the early universe being bright \\citep[i.e., luminosity evolution, preferred by][]{salv07,salv09a,salv09b,pet09} or to an increase in the overall number of GRBs at intermediate and high-$z$ relative to the SFR.  As we discuss below in Section \\ref{sec:rho}, rigorous treatment of the largest available Swift dataset (Section \\ref{sec:redux}) allows for a firm conclusion in favor of rate evolution and not luminosity evolution, and we suggest plausible explanations.  To draw this conclusion and to study GRB rates as a function of intrinsic hardness, flux, and duration (Sections \\ref{sec:fitting} \\& \\ref{sec:energetics}) as well as $z$, we require a detailed model for the Swift satellite detection limit (Section \\ref{sec:limit}). ",
        "conclusions": ""
    },
    "0910/0910.3946_arXiv.txt": {
        "abstract": "We present deep CFHT/MegaCam photometry of the ultra-faint Milky Way satellite galaxies Coma Berenices (ComBer) and Ursa Major II (UMa\\,II). These data extend to $r\\sim25$, corresponding to three magnitudes below the main sequence turn-offs in these galaxies. We robustly calculate a total luminosity of $M_{V}=-3.8\\pm0.6$ for ComBer and $M_{V}=-3.9\\pm0.5$ for UMa\\,II, in agreement with previous results. ComBer shows a fairly regular morphology with no signs of active tidal stripping down to a surface brightness limit of $32.4$ mag arcsec$^{-2}$. Using a maximum likelihood analysis, we calculate the half-light radius of ComBer to be $r_{\\rm half}=74\\pm4$\\,pc ($5.8\\pm0.3\\arcmin$) and its ellipticity $\\epsilon=0.36\\pm0.04$. In contrast, UMa\\,II shows signs of on-going disruption.  We map its morphology down to $\\mu_{V}=32.6$ mag arcsec$^{-2}$ and found that UMa\\,II is larger than previously determined, extending at least $\\sim700$\\,pc ($1.2^\\circ$ on the sky) and it is also quite elongated with an ellipticity of $\\epsilon=0.50\\pm0.2$. However, our estimate for the half-light radius, $123\\pm3$\\,pc ($14.1\\pm0.3\\arcmin$) is similar to previous results. We discuss the implications of these findings in the context of potential indirect dark matter detections and galaxy formation. We conclude that while ComBer appears to be a stable dwarf galaxy, UMa\\,II shows signs of on-going tidal interaction. ",
        "introduction": "For decades, only about a dozen dwarf galaxies were known to orbit the Milky Way.  The majority of these systems corresponded to dwarf spheroidal (dSph) galaxies, the least luminous, but, by number, the dominant galaxy type in the present-day universe. However, over the last five years, and thanks to the advent of the Sloan Digital Sky Survey (SDSS; \\citealt{York2000}) the field of dwarf galaxies in the Milky Way has been revolutionized.  To date, fourteen new systems have been detected as slight overdensities in star count maps using the SDSS data (\\citealt{Willman2005a,Willman2005b}; \\citealt{Belokurov2006,Belokurov2007a,Belokurov2008,Belokurov2009}; \\citealt{Zucker2006a,Zucker2006b}; \\citealt{Sakamoto2006}; \\citealt{Irwin2007}; \\citealt{Walsh2009}). These recent discoveries have revealed a previously unknown population of ``ultra-faint'' systems which have extreme low luminosities, in some cases as low as $L_{V}\\sim300$\\,L$_{\\sun}$ \\citep{Martin2008}, and in average comparable (or lower) to those of Galactic globular clusters. However, spectroscopic surveys of the majority of these systems reveal kinematics and metallicities in line with those of dwarf galaxies (\\citealt{Kleyna2005}; \\citealt{Munoz2006}; \\citealt{SimonGeha2007}; \\citealt{Kirby2008}; \\citealt{Koch2009}; \\citealt{Geha2009}).   Dynamical mass estimates of the ultra-faint galaxies based on line-of-sight radial velocities indicate that these galaxies are extremely dark matter dominated, with central mass-to-light ratios (M/L) as high as $1,000$ in solar units.  These systems are thus good laboratories to constrain cosmological models \\citep[e.g~the 'missing satellites problem';][]{Kauffmann1993,Moore1999,Klypin1999,SimonGeha2007} and to study the properties of dark matter (\\citealt{Strigari2008}; \\citealt{Kuhlen2008}; \\citealt{Geha2009}). However, these applications hinge critically on the assumption that the masses and density distributions in these systems are accurately  known.  Current mass estimates for the ultra-faint dwarfs are based on the assumption that the dynamical state of these systems have not been significantly affected by Galactic tides, and therefore that they are near dynamical equilibrium.   There is circumstantial evidence for past tidal disturbance in a number of these satellites based on morphological studies.  \\citet{Coleman2007} for instance, found the Hercules dSph to be highly elongated, with a major to minor axis ratio of 6:1, larger than for any of the ``classical'' dSphs, and argued for a tidal origin of this elongation. \\citet{Belokurov2007a} reported fairly distorted morphologies for several of the new ultra-faint dwarfs, although \\citet{Martin2008} showed that results based on low number of star may not be statistically significant because they suffer from shot noise. The strongest evidence for tidal interaction perhaps is for Ursa Major II, which appears to be broken into several clumps, lies close to the great circle that includes the Orphan Stream (\\citealt{Zucker2006b}; \\citealt{Fellhauer2007}) and shows a velocity gradient along its major axis \\citep{SimonGeha2007}.  If dSphs galaxies are currently undergoing tidal stripping, or have in the past, then kinematical samples are expected to be ``contaminated'' with unbound or marginally bound stars. This will impact subsequent dynamical modeling, often resulting in overestimated mass contents (\\citealt{Klimentowski2007}; \\citealt{Lokas2009}).  The degree of such contamination is a function of both the projection of the orbit along the line-of-sight and the strength of the tidal interaction.  Therefore, even though tides do not affect appreciably the inner kinematics of dSphs until the latest stages of tidal disruption (e.g., \\citealt{OLA95}; \\citealt{Johnston1999}; \\citealt{Munoz2008}; \\citealt{Penarrubia2008}), studies aimed at identifying the presence of tidal debris can help elucidate the dynamical state of these dwarf galaxies. Obvious tidal features around dSphs would indicate the presence of unbound stars and their effects on the derived masses would need to be investigated.  Alternatively, a lack of clear detections in the plane of the sky would narrow the possibilities that kinematical samples suffer from contamination, although it would not automatically imply that an object has not been tidally affected. Tidal features could still exist but be aligned preferentially along the line-of-sight and thus be hard to detect.  Ultra-faint dwarf galaxies are useful probes of galaxy formation on the smallest scales (\\citealt{Madau2008}; \\citealt{Ricotti2008}).  One of the outstanding questions related to their discovery is whether these galaxies formed intrinsically with such low luminosities, or whether they were born as brighter objects and attained their current luminosities through tidal mass loss. Current metallicity measurements (e.g., \\citealt{Kirby2008}) support the former scenario. They show that the ultra-faint dwarfs are also the most metal poor of the dSphs, following the luminosity-metallicity trend found for the classical dSphs\\footnote{\\citet{SimonGeha2007}, using the same spectroscopic data but a different technique, reported systematically higher metallicites for the ultra-faint dwarfs than those of \\citet{Kirby2008}. The discrepancy is explained by the fact that the former study used the \\citet{Rutledge1997} method that relies on the linear relationship between Ca II triplet equivalent widths and [Fe/H], a technique now known to fail at very low metallicities.}. On the other hand, if tidal features around these objects are firmly detected, it would clearly support the latter hypothesis.  Given the importance of the questions at hand, it is essential that we investigate the dynamical state of these satellites. Due to the extreme low luminosity of the ultra-faint dwarfs, and therefore low number of brighter stars available for spectroscopic studies, deep imaging is currently the only way to efficiently detect the presence of faint morphological features in the outskirts of these systems. We expect any tidal debris to be of very low surface brightness (\\citealt{Bullock2005}) so that detecting it via integrated light is virtually impossible. However, these galaxies are sufficiently nearby to resolve individual stars. Thus, using matched-filter techniques it is possible to detect arbitrarily low surface brightness features ($\\sim 35$ mag arcsec$^{-2}$, \\citealt{Rockosi2002}; \\citealt{Grillmair2006}).  In this article we present the results of a deep, wide-field photometric survey of the Coma Berenices (ComBer) and Ursa Major II (UMa\\,II) dSphs, claimed to be two of the most dark matter dominated dSphs \\citep{Strigari2008}, carried out with the MegaCam imager on the Canada-France-Hawaii Telescope. In \\S2, we present details about the observations and data reduction, as well as of artificial star tests carried out to determine our completeness levels and photometric uncertainties. In \\S3, we recalculate the structural parameters of these systems using a maximum likelihood analysis similar to that of \\citet{Martin2008}. In \\S4 we present the results of our morphological study. We discuss the statistical significance of our density contour maps and the robustness of our result to the variation of different parameters. We show that ComBer looks fairly regular in shape and find no signs of tidal debris down to a surface brightness limit of 32.4 mag arcsec$^{-2}$, whereas UMa\\,II shows a significantly elongated and distorted shape, likely the result of tidal interaction with the Milky Way. Our discussion and conclusions are presented in \\S5 and \\S6, respectively. ",
        "conclusions": "We have carried out a deep, wide-field photometric survey of the Coma Berenices and Ursa Major II dwarf spheroidal galaxies using the MegaCam Imager on CFHT, reaching down to $r\\sim25$ mag, more than three magnitudes below the main sequence turn-offs of these Galactic satellites. This increases roughly by an order of magnitude, with respect to the original SDSS photometry, the number of stars that belong to the respective galaxies and that are available determination of their structural properties and for morphological studies.  Our results can be summarized as follows:  1. We used a maximum likelihood analysis similar to the one used by \\citet{Martin2008} and \\citet{Sand2009} to calculate structural parameters for three different density profiles: King, Plummer and exponential. We find characteristic sizes of $r_{\\rm half}=74\\pm4$\\,pc ($5.8\\pm0.3\\arcmin$) and $123\\pm3$\\,pc ($14.1\\pm0.3\\arcmin$) for ComBer and UMa\\,II respectively (from the exponential profile). Our results provide much tighter constraints on these structural parameters than possible with previous datasets, but are consistent with earlier determinations using SDSS photometry.  2. We have re-calculated the total luminosities for both systems and find, for ComBer $M_{V}=-3.8\\pm0.6$ and for UMa\\,II $M_{V}=-3.9\\pm0.5$ (for a choice of Salpeter IMF), in very good agreement with previous results, confirming that ComBer and UMa\\,II are among the faintest of the known dwarf satellites of the Milky Way. We have also used a \\citet{Chabrier2001} IMF but the results remain virtually unchanged.  3. We have found that ComBer shows a fairly regular morphology with no clear detection of potential stripped material down to a surface brightness of $32.4$ mag arcsec$^{-2}$. Additionally, its number density profile is reasonably well matched by a choice of either King, Plummer or exponential profile. We thus conclude that ComBer is likely a stable dwarf galaxy which would make it one of the most dark matter dominated of the dSph systems.  4. We have also studied the morphology of UMa\\,II and find that, unlike ComBer, it shows signs of being significantly disrupted. UMa\\,II is larger than previously determined, extending at least $\\sim700$\\,pc ($1.2^\\circ$ on the sky) and it is also quite elongated. Its density profile and overall shape resemble a structure possibly in the latest stages of tidal destruction. Furthermore, its number density profile is not well matched by neither of the three profiles we tried and it is much better described by two power laws, further supporting a tidal scenario.  5. The overall 2D surface density distributions of both systems are not affected by shot-noise and are therefore robust. We find no evidence for isolated tidal debris beyond the main bodies of ComBer and UMa\\,II to our surface brightness limits of $32.4$ and $34.8$ mag arcsec$^{-2}$, respectively.  Deep, wide-field imaging of the recently discovered ultra-faint galaxies currently lags behind spectroscopic observations of these objects. We show in this paper that high quality, deep photometry is an equally important tool in studying the dynamical state of the ultra-faint dwarfs. These data can also be used to constrain the star formation histories of ComBer and UMa\\,II which we will explore in a future contribution.  We acknowledge David Sand, Josh Simon and Gail Gutowski for useful discussions and Peter Stetson for graciously providing copies of DAOPHOT and ALLFRAME. M.G and B.W. acknowledge support from the National Science Foundation under award number AST-0908752."
    },
    "0910/0910.3207_arXiv.txt": {
        "abstract": "We lay out the framework to numerically study nonlinear structure formation in the context of scalar-field-coupled cold dark matter models ($\\varphi$CDM models) where the scalar field $\\varphi$ serves as dynamical dark energy. Adopting parameters for the scalar field which leave negligible effects on the CMB spectrum, we generate the initial conditions for our $N$-body simulations. The simulations follow the spatial distributions of dark matter and the scalar field, solving their equations of motion using a multilevel adaptive grid technique. We show that the spatial configuration of the scalar field depends sensitively on the local density field. The $\\varphi$CDM model differs from standard $\\Lambda$CDM at small scales with observable modifications of, e.g., the mass function of halos as well as the matter power spectrum.  Nevertheless, the predictions of both models for the Hubble expansion and the CMB spectrum are virtually indistinguishable. Hence, galaxy cluster counts and weak lensing observations, which probe structure formation at small scales, are needed to falsify this class of models. ",
        "introduction": "The origin and nature of dark energy (Copeland~et~al.~2006) is one of the most difficult challenges facing physicists and cosmologists at the present time. Among all the proposed models to tackle this problem, the introduction of a scalar field is perhaps the most popular.  The scalar field, denoted by $\\varphi$, should have no coupling to normal matter to be consistent with stringent constraints from experiments (Will 2006, and references therein), but could couple to the dark matter, therefore producing a fifth force between dark matter particles. This idea has gained a lot of interest in recent years because dark matter physics are unknown, and such a coupling could alleviate the coincidence problem of dark energy (e.g., Amendola 2000; Chiba 2001; Chimento et al. 2003). Furthermore, it is commonly predicted by low energy effective theories derived from a more fundamental theory. A specific and interesting possibility is the chameleon mechanism (Khoury \\& Weltman 2004; Mota \\& Shaw 2006), by virtue of which the scalar field acquires a large mass in high density regions and thus the fifth force becomes undetectable on short ranges, thus also evading constraints from the large-scale cosmic microwave background (CMB). Indeed, at the linear perturbation level, there have been a lot of studies about the coupled scalar field and $f(R)$ gravity models (e.~g., Li \\& Barrow 2007; Hu \\& Sawicki 2007).  Nevertheless, little is known about these models on nonlinear scales. It is well known that the matter distribution at late times, i.e. $z\\lesssim 2$ for cluster scales, evolves in a nonlinear way, making the behavior of the scalar field more complex and the linear analysis insufficient to produce accurate results that can be confronted with observations. For the latter purpose, the best way forward is to perform full $N$-body simulations (Bertschinger 1998) to evolve the individual particles step by step.  $N$-body simulations including scalar fields and related models have been performed before (Linder \\& Jenkins 2003; Mainini~et~al.  2003; Macci\\`o~et.~al. 2004; Springel \\& Farrar 2007; Kesden \\& Kamionkowski 2006a, 2006b; Farrar \\& Rosen 2007; Baldi~et~al.  2008; Oyaizu 2008; Keselman~et~al. 2009; Li \\& Zhao 2009). For example, in the work of Macci\\`o~et~al.~(2004) the simulations included several effects due to the coupling between dark energy and dark matter (e.g. modified gravitational constant, an extra dragging term in Newton's equations and time variable dark matter particle masses), but did not consider a spatial variation of the dark energy scalar field.  The more complete simulation of the scalar field by Li \\& Zhao (2009) shows that this approximation is only good for a limited choice of parameters and the scalar field potential. Here we extend the work of Li \\& Zhao (2009).  This paper is organized as follows: In \\S~\\ref{sect:description}, we shall briefly review the general equations of motion for the coupled scalar field model introduced in Li \\& Zhao (2009), and present our specific choices of the coupling function and the scalar field potential. In \\S~\\ref{sect:non-linear}, we describe the formulae and the algorithm of the $N$-body simulation, analyze the results of our coupled scalar field $N$-body simulations, compare it with that of the standard $\\Lambda$CDM model, and explain the physical origin of the new features. Finally, we conclude and discuss observational implications in \\S~\\ref{sect:disc}. ",
        "conclusions": "\\label{sect:disc}  We have presented a general framework to study nonlinear structure formation in coupled scalar field models, in particular the models of Li \\& Zhao (2009). While these models are virtually indistinguishable from $\\Lambda$CDM on both very large and very small scales, intermediate scales at low redshift ($z\\lesssim 1$) relevant for galaxy clusters ($\\sim 10^{2}-10^{3}$kpc) open a new window to test and constrain the interesting part of the parameter space.  On these scales, the matter power spectrum is significantly increased compared to that of $\\Lambda$CDM. Observationally, this would most likely appear as a change of $\\sigma_{8}$ on the order of $15$-$20$\\% for models with $\\gamma=0.5-1$ and $\\mu=10^{-6}$ (see Fig.~2). Any variation of $\\sigma_8$ seems to be lower than 30\\% for current lensing measurements such as the CFHT Legacy Survey (e.g., Hoekstra et al. 2006; Fig.~11 of Fu et al. 2008) over a rather limited range; however, future surveys, such as the Kilo-Degree Survey (KIDS), will be able to measure the scale dependence within the range $k=0.1-10h$Mpc$^{-1}$, where the deviation of the models from $\\Lambda$CDM is maximal."
    },
    "0910/0910.4611_arXiv.txt": {
        "abstract": "We present an improved analysis of halo substructure traced by RR Lyrae stars in the SDSS stripe 82 region. With the addition of SDSS-II data, a revised selection method based on new $ugriz$ light curve templates results in a sample of 483 RR Lyrae stars that is essentially free of contamination. The main result from our first study persists: the spatial distribution of halo stars at galactocentric distances 5--100 kpc is highly inhomogeneous. At least 20\\% of halo stars within 30 kpc from the Galactic center can be statistically associated with substructure. We present strong direct evidence, based on both RR Lyrae stars and main sequence stars, that the halo stellar number density profile significantly steepens beyond a Galactocentric distance of $\\sim$30 kpc, and a larger fraction of the stars are associated with substructure. By using a novel method that simultaneously combines data for RR Lyrae and main sequence stars, and using photometric metallicity estimates for main sequence stars derived from deep co-added $u$-band data, we measure the metallicity of the Sagittarius dSph tidal stream (trailing arm) towards R.A.$\\sim2^h-3^h$ and $Dec\\sim0^\\circ$ to be 0.3 dex higher ($[Fe/H]=-1.2$) than that of surrounding halo field stars. Together with a similar result for another major halo substructure, the Monoceros stream, these results support theoretical predictions that an early forming, smooth inner halo, is metal poor compared to high surface brightness material that have been accreted onto a later-forming outer halo. The mean metallicity of stars in the outer halo that are not associated with detectable clumps may still be more metal-poor than the bulk of inner-halo stars, as has been argued from other data sets. ",
        "introduction": "}  Studies of the Galactic halo can help constrain the formation history of the Milky Way and the galaxy formation process in general. For example, within the framework of hierarchical galaxy formation \\citep{fbh02}, the spheroidal component of the luminous matter should reveal substructures such as tidal tails and streams \\citep{jhb96,hw99,bkw01,har01}. The number of these substructures, due to mergers and accretion over the Galaxy's lifetime, may provide a crucial test for proposed solutions to the ``missing satellite'' problem \\citep{bkw01}. Substructures are expected to be ubiquitous in the outer halo (galactocentric distance $>15-20$ kpc), where the dynamical timescales are sufficiently long for them to remain spatially coherent \\citep{jhb96,may02}, and indeed many have been discovered \\citep[e.g.,][]{iba97,yan00,ive00,bel06,gri06,vz06,bel07a,new07,jur08}. Various luminous tracers, such as main sequence turn-off stars, RR Lyrae variables, or red giants are used to detect halo substructures, and of them, RR Lyrae stars have proven to be especially useful.  RR Lyrae stars represent a fair sample of the old halo population \\citep{smi95}. They are nearly standard candles ($\\langle M_V \\rangle = 0.59 \\pm 0.03$ at $[Fe/H]=-1.5$, \\citealt{cc03}), are sufficiently bright to be detected at large distances ($5-120$ kpc for $14 < V < 21$), and are sufficiently numerous to trace the halo substructure with good spatial resolution ($\\sim1.5$ kpc at 30 kpc). Fairly complete and relatively clean samples of RR Lyrae stars can be selected using single-epoch colors \\citep{ive05}, and if multi-epoch data exist, using variability (\\citealt{ive00}; QUEST, \\citealt{viv01}; \\citealt{ses07}, \\citealt{dle08}; SEKBO, \\citealt{kel08}; LONEOS-I, \\citealt{mic08}). A useful comparison of recent RR Lyrae surveys in terms of their sky coverage, distance limits, and sample size is presented by \\citet{kel08} (see their Table 1). Compared to other surveys, the so-called stripe 82 survey (an equatorial strip defined by declination limits of $\\pm$1.27$^\\circ$ and extending from R.A.$\\sim$20$^h$ to R.A.$\\sim$4$^h$) based on Sloan Digital Sky Survey data (SDSS; \\citealt{yor00}) is the deepest one yet obtained (probing distances to $\\sim110$ kpc), and the only survey with available 5-band photometry.  In an initial study based on SDSS stripe 82 data, we discovered a complex spatial distribution of halo RR Lyrae stars at distances ranging from 5 kpc to $\\sim$100 kpc \\citep[hereafter S07]{ses07}. The candidate RR Lyrae sample was selected using colors and low-order light curve statistics. In this paper, we present a full light curve analysis enabled by the addition of new stripe 82 observations. The resulting sample of RR Lyrae stars is highly complete and essentially free of contamination. The extended observations now allow the determination of flux-averaged magnitudes, thus providing better luminosity and distance estimates. This reduces the uncertainty in the spatial distribution of RR Lyrae stars, and increases the overall accuracy of computed density maps.  While this paper was in preparation, another study of RR Lyrae stars in stripe 82, based on essentially the same data set, was announced \\citep{wat09}. In addition to analysis that overlaps with their work, here we also discuss \\begin{itemize} \\item An RR Lyrae template construction for the SDSS bandpasses, and analysis of the template behavior \\item Template fitting and visual inspection of best-fits for single-band and color light curves, which results in an exceedingly clean final sample \\item Extended color range for the initial selection, which results in about 20\\% more stars than in the sample analyzed by Watkins et al. \\item A step-by-step analysis of the sample incompleteness at the faint end \\item A new method for constraining the metallicity of spatially coherent structures that are detected using main sequence and RR Lyrae stars \\item A comparison of the observed RR Lyrae spatial distribution with predictions based on an oblate halo model constrained by SDSS observations of main sequence stars. \\end{itemize} We emphasize, however, that both studies obtain consistent main results: strong substructure in the spatial distribution of halo RR Lyrae stars and evidence for a steepening of the best-fit power-law description beyond a galactocentric radius of $\\sim$30 kpc.  Our data set and the initial selection of candidate RR Lyrae stars are described in \\S~\\ref{data}, while the construction of light curve templates is presented in \\S~\\ref{ugriz_templates}. Various properties of the resulting RR Lyrae sample, with emphasis on the spatial distribution, are analyzed in \\S~\\ref{sec:spatial}. In \\S~\\ref{sec:MS} we compare the spatial distribution of RR Lyrae stars and main sequence stars and constrain the metallicity distribution of the Sgr tidal stream. Our main results are summarized and discussed in \\S~\\ref{discussion}. ",
        "conclusions": "}  We have presented the most complete sample of RR Lyrae stars identified so far in the SDSS stripe 82 data set, consisting of 379 RRab and 104 RRc stars. Our visual inspection of single-band and color light curves insures that the sample contamination is essentially negligible. Although the sky area is relatively small compared to other recent surveys, such as \\citet{kel08} and \\citet{mic08}, this RR Lyrae sample has the largest distance limit to date ($\\sim$100 kpc).  In addition, a large number of well-sampled light curves in the $ugriz$ photometric system have enabled the construction of a new set of empirical templates. This multi-band empirical template set provides strong constraints for models of stellar pulsations, through the single-band light curve shape, the shape variation with wavelength, and through the distributions of the templates. For example, while $ab$-type RR Lyrae appear to have a continuous distribution of light curve shapes, $c$-type RR Lyrae exhibit a clear bimodal distribution. This bimodality suggests that the latter class may include two distinct subclasses, which would be an unexpected result given that RR Lyrae stars have been carefully studied for over a century (e.g., \\citealt{bai03}). A subclass of RR Lyrae stars with amplitudes and periods similar to the $c$-type, but with asymmetric light curves, has been noted in the MACHO data by \\citet{alc96}. They suggested that these stars are pulsating in the second-harmonic mode.  Our accurate templates will greatly assist the identification of RR Lyrae stars in data sets with relatively small numbers of epochs. For example, all three major upcoming wide-area ground-based surveys plan to obtain fewer than 10 epochs (SkyMapper, Dark Energy Survey and Pan-STARRS1 3$\\pi$ survey), in the same photometric system. Without template fitting, the selection based on only low-order light-curve statistics includes a significant fraction ($\\sim$30\\%) of contaminants (dominated by $\\delta$ Scuti stars).  The high level of completeness and low contamination of our resulting sample, as well as the more precise distance estimates, enabled a more robust study of halo substructure than was possible in our first study. Our main result remains: the spatial distribution of halo RR Lyrae at galactocentric distances 5--100 kpc is highly inhomogeneous. The distribution of $\\rho/\\rho^{RR}_{model}$ suggests that at least 20\\% of the halo within 30 kpc is found in apparently real substructures. \\citet{bel08} find a lower limit of 40\\%, but they used a different metric and different tracers (turn-off stars). Schlaufman et al.~(in prep) have demonstrated that at least 34\\% of inner-halo (in their paper they explored regions up to 17 kpc from the Galactic center) main sequence turnoff stars are contained in structures they refer to as ECHOS (elements of cold halo substructures), based on an examination of radial velocities in the SEGUE sample. Thus, three separate studies, using very different probes, have come to essentially identical conclusions.  A comparison of the observed spatial distribution of RR Lyrae stars and main sequence stars to the \\citet{jur08} model, which was constrained by main sequence stars at distances up to 20 kpc, strongly suggests that the halo stellar number density profile steepens beyond $\\sim$30 kpc. While various indirect evidence for this behavior, based on kinematics of field stars, globular clusters, and other tracers has been published (\\citealt{car07}; and references therein), our samples provide a direct measurement of the stellar halo spatial profile beyond the galactocentric distance limit of $\\sim$30 kpc. A similar steepening of the spatial profile was detected using candidate RR Lyrae stars by \\citet{kel08}. In addition to presenting new evidence for the steepening of halo stellar number density profile beyond $\\sim$30 kpc, we have confirmed the result of \\citet{mic08} that the density profile for Oosterhof II stars is steeper than for Oosterhof I stars within 30 kpc.  We have analyzed several methods for estimating the metallicities of RR Lyrae stars. Using spectroscopic data processed with the SDSS SSPP pipeline as a training sample, we have established a weak correlation with the $u-g$ color at minimum light. The rms scatter around the best-fit relation is $\\sim0.3$ dex, but systematic errors could be as large, and the best-fit relation should be used with care. We have established a relationship between the parametrization of our template set for $ab$-type RR Lyrae stars and the results based on the Fourier expansion of light curves. This allowed us to estimate the metallicity from observed light curves with an estimated rms scatter of $\\sim0.3$ dex. This method failed to uncover a systematic metallicity difference between field RR Lyrae stars and those associated with the Sgr stream, which is seen in main sequence stars. Nevertheless, it is likely that this is a promising method for estimating metallicity, which is currently limited by the lack of a large and reliable calibration sample. Our templates can be used to bypass noisy Fourier transformation of sparsely sampled data, as already discussed by \\citet{kin06} and \\citet{kk07}.  We introduced a novel method for estimating metallicity that is based on the absolute magnitude vs.~metallicity relations for RR Lyrae stars and main sequence stars (calibrated using globular clusters). This method does not require the $u$-band photometry, and will be useful to estimate metallicity of spatially coherent structures that may be discovered by the Dark Energy and Pan-STARRS surveys. While the existing SDSS data are too shallow to apply this method to the Pisces stream, we used it to obtain a metallicity estimate for the Sgr tidal stream that is consistent with an independent estimate based on the photometric $u$-band method for main sequence stars. Our result, $[Fe/H]=-1.2$, with an uncertainty of $\\sim$0.1 dex, strongly rules out the hypothesis that the trailing arm of the Sgr dSph tidal stream has the same metallicity as halo field stars ($[Fe/H]=-1.5$), as suggested by \\citet{wat09}. \\citet{yan09} detected peaks at $[Fe/H]=-1.3$ and at $[Fe/H]=-1.6$ using SDSS spectroscopic metallicities for blue horizontal-branch (BHB) stars from the trailing arm. However, both the Yanny et al.~and Watkins et al.~results could be affected by systematic errors in the SDSS metallicity scale for BHB stars. Our results suggest that the SDSS metallicity scale for BHB stars could be biased low by about 0.3 dex (relative to the SDSS metallicity scale for main sequence stars, and assuming that RR Lyrae and main sequence stars from the Sgr tidal stream in stripe 82 area have the same metallicity distributions).  Simulations by \\citet{bj05} and \\citet{joh08} predict that there should be a difference in the chemical composition between stars in the inner halo that was built from accretion of massive satellites about 9 Gyr ago, and outer halo dominated by stars coming from dSph satellites that were accreted in the last 5 Gyr (see Fig.~11 in \\citealt{bj05}). These accreted dSph satellites were presumably more metal-poor than the massive satellites accreted in earlier epochs (see Fig.~3 in \\citealt{rob05}). Other recent simulation studies support these conclusions. For example, \\citet{dh08} find evidence in their simulation for a strong concentration of (relatively) higher metallicity stars at distances close to the Galactic center, and the presence of (relatively) lower metallicity stars at distances beyond 20 kpc from the center. \\citet{zol09} find from their simulations that their inner halos include stars from both in-situ formed stars and accreted populations, while their outer halos appear to originate through pure accretion and disruption of satellites. Salvadori et al.~(in prep) have considered the distribution of ages and metallicities of metal-poor stars in a Milky-Way like halo, as a function of galactocentric radius, based on a hybrid N-body and semi-analytic simulation. They find an inner-halo population that is well-described by a power-law index $n=-3.2$ (for stars with $-2<[Fe/H]<-1$), and an outer-halo consistent with a much shallower profile, $n=-2.2$. The relative contributions of stars with $[Fe/H]\\le-2$ in their simulation increases from about 16\\% for stars within 7 kpc $< R <$ 20 kpc, to $>40\\%$ for stars with $R >$ 20 kpc.  Our observational results provide some support for these simulation-based predictions. The faintest main sequence stars ($r\\sim21.5$), not including those apparently associated with the Sgr stream, exhibit median metallicities at least 0.3 dex lower ($[Fe/H]=-1.8$, see the left panel in Figure~\\ref{fig:bsFig3}) than halo stars within 10-15 kpc from the Galactic center \\citep{ive08}. Other studies provide additional support. For example, \\citet{car07} and Carollo et al.~(in prep) have indicated that field stars likely to be associated with the outer halo exhibit metallicities that are substantially lower ($[Fe/H]=-2.2$) than those of the inner halo (they place the inner/outer halo boundary at $\\sim$10 kpc). The two metallicity measurements are fully consistent because photometric $u$-band metallicity estimator is biased high for spectroscopic values with $[Fe/H]< -2$ (Bond et al., in prep).  Carollo et al.~(in prep) also found that the inferred density profiles of the inner- and outer-halo populations differ as well; the inner halo being consistent with a power-law profile with index $n \\sim -3.4$, and the outer halo having index $n = -2.1$ (the data analysed by Carollo et al. are not suitable for determination of the relative normalizations of the inner- and outer-halo populations, due to the selection effects involved). This difference in the density profiles is very similar to the difference in density profiles for Oosterhof I and Oosterhof II stars. However, we emphasize that the steeper profile for the inner halo from Carollo et al.~corresponds to more metal-rich stars, while we found weak evidence that Oosterhof II stars tend to be more metal poor than Oosterhof I stars.  Simulations indicate that high surface brightness substructures in the halo originate from single satellites, typically massive dSph which tend to be accreted over the last few Gyr \\citep{bj05}, and these massive galaxies are expected to be more more metal-rich than halo field stars \\citep{fon08}. The results from \\citet{ive08} and the results presented here seem to support this prediction. The inner halo has a median metallicity of $[Fe/H]=-1.5$, while at least two strong overdensities have higher metallicities -- the Monoceros stream has $[Fe/H]=-1.0$, and for the trailing part of the Sgr tidal stream we find $[Fe/H]=-1.2$. We emphasize that these three measurements are obtained using the same method/calibration and the same data set, and thus the measurements of relative differences are expected to be robust. Simulations by \\citet{joh08} predict that the metallicity of low surface brightness features, such as the Virgo Overdensity \\citep{jur08}, is expected to be lower than the median halo metallicity. Although the initial estimate by \\citet{jur08} (see their Figure 39) claimed a metallicity for this structure of $[Fe/H]\\sim-1.5$, a recent study of the photometric metallicity of stars in this structure by \\citet{an09} have suggested that the mean metallicity is even lower: $[Fe/H]= -2.0$. Both of these studies are limited by potentially large systematics associated with the photometric metallicity technique, so more detailed analysis of spectroscopically determined metallicities for stars in this structure would be of great value. The sparse information that is available, based on spectroscopic determinations, supports a mean metallicity between $[Fe/H]= -1.8$ and -2.2 \\citep{duf06,viv08,pri09,cas09}. \\citet{cas09} report the measurement of a precise absolute proper motion for the RR Lyrae star RR 167, which appears highly likely to be associated with the Virgo Overdensity. This proper motion, in combination with their distance estimate (17 kpc from the Galactic center) and radial velocity measurement, indicate that the Virgo Overdensity structure may well be on a highly destructive orbit, with pericenter $\\sim 11$ kpc, and apocenter $\\sim 90$ kpc. Thus, the interpretation of the Virgo Overdensity as a dwarf galaxy, perhaps similar to those that participated in formation of the outer-halo population, is strengthened.  Our result that (inner-) halo stellar number density profile steepens beyond $\\sim$30 kpc is limited by the relatively small distance limit for main sequence stars (35 kpc), the sparseness of the RR Lyrae sample ($\\sim$500 objects), and the small survey area ($\\sim$300 deg$^2$). Ideally, the halo stellar number density profile should be studied using numerous main sequence stars detected over a large fraction of sky. To do so to a distance limit of 100 kpc, imaging in at least $g$ and $r$ bands (or their equivalent) to a depth several magnitudes fainter than the co-added SDSS stripe 82 data is required ($r>25$). Pan-STARRS, the Dark Energy Survey and LSST are planning to obtain such data over large areas of sky. The LSST, with its deep $u$-band data, will also extend metallicity mapping of field main sequence stars over half of the sky in the south; see \\citet{ive08} for details. For substructures to be potentially discovered in the north by Pan-STARRS, the method presented here can be used to estimate the metallicity of spatially coherent structures even without the $u$-band data."
    },
    "0910/0910.1728_arXiv.txt": {
        "abstract": "The correlation between the large-scale distribution of galaxies and their spectroscopic properties at $z=1.5$ is investigated using the Horizon MareNostrum cosmological run.  We have extracted a large sample of $10^5$ galaxies from this large hydrodynamical simulation featuring standard galaxy formation physics. Spectral synthesis is applied to these single stellar populations to generate spectra and colours for all galaxies. We use the skeleton as a tracer of the cosmic web and study how our galaxy catalogue depends on the distance to the skeleton. We show that galaxies closer to the skeleton tend to be redder, but that the effect is mostly due to the proximity of large haloes at the nodes of the skeleton, rather than the filaments themselves.  This effects translate into a bimodality in the colour distribution of our sample. The origin of this bimodality is investigated and seems to follow from the ram pressure stripping of satellite galaxies within the more massive clusters of the simulation.  The virtual catalogues (spectroscopical properties of the MareNostrum galaxies at various redshifts) are available online at {\\tt http://www.iap.fr/users/pichon/} {\\tt MareNostrum/catalogues}. ",
        "introduction": "\\label{sec:intro}   \\begin{figure} \\includegraphics[width= 1.1\\columnwidth]{fig/blah} \\caption{A 3D view of the galaxies and the skeleton of the dark matter of  MareNostrum at $z=1.6$. The box is 50 $h^{-1}$Mpc aside.} \\label{fig:skel} \\end{figure}  During the past decade, the $\\Lambda {\\rm {CDM}}$ cosmological model of the Universe has been established as the framework of choice in which to interpret how and when observed galaxies acquire their properties. Arguably the most important feature of this framework is to provide us with an explanation as to why  many  of  these properties (physical sizes, luminosities) strongly correlate with galaxy mass while others (star formation rates, morphological type) do not seem to. Unsurprizingly, the all-time favored culprit is the interplay between galaxies and the intergalactic medium (IGM) at large. In other words, the large scale environment of galaxies is claimed to play an important role in shaping some of their properties, while the rest of them are thought to depend solely on small scale (internal) processes. However, having said that, one still has to determine for which of these properties ``Nurture\" dominates over ``Nature\" and therein lies the whole difficulty of the issue.  Indeed, since the early 70s, there has been a plethora of studies devoted to measuring the impact of environment on galaxy properties. \\cite{DavisGeller1976} first pointed out that early-type galaxies are more strongly clustered than the late types, while \\cite{Dressler1980} and \\cite{PostmanGeller1984} demonstrated the existence of a morphology-density relation (MDR). Following in their footsteps, \\cite{Balogh1998}, \\cite{Hashimoto} and more recently \\cite{Christlein} and \\cite{Poggianti} systematically showed that galaxies living in denser environments tend to be redder and have lower star formation rates (SFRs) than their more isolated counterparts. One can think of several physical processes associated with different types of environment that could play a role in causing such alterations. More specifically, they include, by ascending order of environment density: (i) major (wet) mergers, which can turn spiral galaxies into ellipticals (e.g. \\cite{toomre}), and drive a massive starburst wind which quenches future star formation by ejecting the interstellar medium (ISM) out of galaxies (\\cite{Mihos,maclow}); (ii) active galactic nuclei (AGN) or shock-driven winds (\\cite{norm,sprin}; (iii) galaxy ``harassment\" (rapid encounters or fly-bys which dominate over mergers in rich clusters) causing discs to heat and possibly triggerring the build-up of a bulge via the formation of a bar (\\emph{e.g.} \\cite{moore}). For the latter category, the diffuse gas associated with the galaxies' host dark matter (sub)halo which constitutes the main fuel supply for future star formation can also be stripped, thus suppressing later star formation by ``starvation\" or ``strangulation\" \\citep{larson,Bekki}. Moreover, part of their ISM can also be pulled out of these galaxies, either by tidal forces arising from the gravitational potential of the cluster or by ram pressure stripping by the intracluster medium (ICM) \\cite{Gunn,Abadi,Chung}.  \\\\ Although these latter environment dependent processes seem potent enough, recent work carried out by  \\cite{Tanaka} and \\cite{vdB}  indicate that they might not be the main mechanisms for quenching star formation activity. This claim is corroborated by the higher redshift results ($ z \\!\\sim\\! 1$) obtained with the DEEP2 (e.g. \\cite{Cooper2006,Cooper2007,Cooper2008,Gerke2007,Coil2008}), (z--)COSMOS \\citep{Scoville2007,Cassata2007,Tasca2009} and VVDS \\citep{Scodeggio} surveys where clusters are more scarce, along with the fact that morphological and spectrophotometric properties of local galaxies are also found to be correlated with their internal properties, such as luminosity, mass or internal velocity (e.g. \\cite{Kauffmann2003}). Here as well, one can invoke various physical processes to explain such dependences on internal properties. Supernova feedback can heat and stir the ISM, possibly ejecting large amounts of gas out of galaxies and it is expected to scale with galaxy mass \\citep{Larson1974,Dekel1986}. There also exists a growing host of observational evidence that AGN play a key part in quenching star formation \\citep{Schawinski2006,Schawinski2007,Salim2007} and this AGN feedback should likely impact more massive galaxies since they host larger mass black holes \\citep{Magorrian,Silk1998}. In light of these investigations, it becomes apparent that disentangling nature and nurture is a more complicated process than one would naively have thought to begin with.   Clearly, the fundamental requirement to tackle this issue is to properly characterize the anisotropic  environment of galaxies, both in observational samples and theoretical models, spanning as broad a range of environments as possible, from isolated field galaxies to groups and rich clusters. The vast majority of the studies in the literature accomplish this task either by counting the number of neighbours that a galaxy has within a fixed aperture on the sky or by measuring the distance to the $n^{\\rm th}$ nearest galaxy, where $n$ is an integer in the range 3\u201310. Although these indicators are straightforward to obtain, their physical interpretation, let alone their comparison to theoretical models are far from being straightforward (\\emph{cf.} \\cite{Kauffmann2004,Weinmann2006}). Meanwhile, looking at the distribution of observed galaxies in modern cosmological surveys, such as the 2dF \\citep{2dF} and the SDSS \\citep{SDSS}, the most striking feature is that they look organised along linear structures linking clusters together (see Figure \\ref{fig:skel}). This filamentary network, dubbed as the ``Cosmic Web\" \\citep{bond}, has a dynamical origin and reflects the anisotropic accretion taking place in clusters \\citep{skel1}. It therefore seems natural to describe the environment of galaxies in terms of their location with respect to these filaments in order to investigate the influence of the Cosmic Web on the properties of the galaxies it encompasses.   In this  paper, we carry out such a study at intermediate redshift ($z\\sim1.5$), mainly from a theoretical perspective, using a recent diagnostic tool to characterize the 3D environment called the skeleton \\citep{sousb08}, which we combine with the largest  hydrodynamical cosmological simulation performed to date \\citep{ocvirk2008,prunet2008, dekelnature}. Run on the MareNostrum computer at the Barcelona Supercomputer Center using  the RAMSES  code \\citep{teyssier02}, this simulation is one of the flagship simulations realized by the  Horizon collaboration ({\\tt http://www.projet-horizon.fr}). It includes   a detailed treatment of metal--dependent gas cooling, UV heating, star formation, supernovae feedback and metal enrichment.  Specifically, we  will address the question: are the physical conditions \\emph{within the filaments} dramatic enough to strongly  influence the properties of the galaxies it encompasses?   The  outline  of  this  paper  is  as  follows:  first  we  describe  in Section \\ref{sec:method} our methodology,   in  terms   of  numerical   techniques, estimators  and  statistical measurements.   The dependence of the spectroscopic properties on the filamentary environment is then  discussed in Section \\ref{sec:fil}, while Section \\ref{sec:bimodality} investigates the observed bimodality and discusses comparison to observations, and Section \\ref{sec:conclusion} wraps up.   Some checks are performed in Appendix A, Appendix B describes the  publicly available catalogues and Appendix C sums up the subgrid physics used. ",
        "conclusions": "\\label{sec:conclusion}   The Cosmic Web is a key feature of the organisation of galaxies on large scales. This paper investigated the influence of this filamentary environment on the spectroscopic properties of galaxies, using the MareNostrum simulation which was postprocessed using stellar population synthesis. The cosmic web was traced  with the skeleton algorithm, and  has proven here to be a very effective mean of probing the anisotropy of the large-scale structures.  We found gradients of spectroscopic properties of galaxies with the distance to filaments, but demonstrated that they can be explained by the fact that the distance to filaments is biased by the distance to the nodes of the network (group or clusters of galaxies). Two procedures were introduced to remove the influence of these nodes: (i) volume-averaging  this distance decreases the influence of the compact, dense regions and focus on wider structures, such as filaments; (ii) pair comparison which seeks the filamentary counterpart of each galaxy. Both methods show that the influence of the filaments alone is negligible compared to influence of clusters. A bimodality in colour was also found to occur below redshift $z\\approx 2$ and its origin was investigated. It  is due to a population of red, small galaxies ($\\sim 10^8 M_\\odot$) accreting on the nodes of the Cosmic Web, while more massive objects, $M> 10^8 M_\\odot$, are mostly unaffected. These galaxies have their star formation quenched while they approach the clusters. It remains to be confirmed that this stripping process is not amplified by a lack-of-resolution effect, since (i) it involves  amongst the smallest (virtual) galaxies in the simulation, and (ii) it seem to create a tension with observations at redshift zero \\citep{kimm}.  These findings suggest that the large-scale filaments are only dynamical features of the density field, reflecting the flow of galaxies accreting on clusters; the conditions in the filaments are not dramatic enough to influence strongly the properties of the galaxies it encompases, unlike the intra (proto) cluster medium, which seems able to strip down the ISM of the  incoming low mass galaxies. Appendix \\ref{sec:test_cut} shows that  this statement remains valid when the study is limited to the most important filaments.   Finally, we did  find a weak metallicity gradient away from the filaments which could reflect the large-scale inhomogeneity in the distribution of metals (the so called WHIM).   One could have imagined that even if the  large-scale filamentary network were  purely a tracer of the large-scale dynamics, galaxies within the large-scale filaments should be redder and older, since they would have joined the cosmic super highway earlier on average when compared to field galaxies. This effect is not seen at those redshifts, as (i) older galaxies continue to accrete new cold gas on small scales and form stars, hence remain blue, (ii) some field galaxies continuously join the large filamentary network, and (iii) on large-scales, a significant fraction of the matter is accreted more or less radially onto the nodes, while filaments are  collecting  left-overs from this accretion more indirectly \\citep{skel1}.  Recently \\citep{keres05,ocvirk2008,dekelnature},  it was emphasized in steps that anisotropic metal-rich cold stream accretion regulates the inflow of cold gas towards the inner regions of the most massive galaxies of the high redshift universe ($z>2$). It was then conjectured that this process could explain  the observed  bimodality of spectroscopic galactic properties at lower redshift.  In this paper, we have shown that the geometric distribution of the colour of galaxies is not sensitive to the detailed large-scale filamentary network, but only to its nodes. This apparent paradox may be lifted when noting that the self-regulating anisotropic filamentary accretion  occurs on much smaller scales, and was quantified for the central      galaxies of the simulation which {\\sl are} sitting at the nodes of the network; in contrast, when considering the full galactic population, a typically low mass galaxy is not transformed by its encounter with the different physical condition of the weakly overdense intergalactic medium within the cosmic web, unless it falls into the intra cluster medium of a large node.\\\\ In other words, massive galaxies feel the small scale filaments (cold streams) feeding them at nodes; low mass galaxies are not spectroscopically changed while entering the large-scale filaments. The mesoscopic (below a Mpc scale) filamentary feature of the cosmic network may geometrically solve the self-regulating process of galactic accretion, but  we have demonstrated here that its large-scale counterpart does not seem to directly affect the colours of galaxies.  In \\cite{sousb08}, the effect of redshift distortion is partially addressed, and the corresponding  algorithm is now being extended to discrete surveys via a Delaunay tessellation (which are therefore not sensitive to edge effects). This, as argued in Section~\\ref{sec:discussion} is a critical step towards performing a similar 3D analysis on real data, which as we have shown is essential to quantify these gradients, as in projection, the information is lost (see Figure~\\ref{fig:2D}). In particular, it would  be of great interest to carry out these measurements on a  DEEP-2/VVDS/z-COSMOS-like survey as well as to bring the simulation down to lower redshift and reach a time in cosmic history when upcoming large  observational surveys (\\emph{e.g.} LSST \\cite{LSST}, BOSS \\cite{BOSS}) overlap statistically with the predictions of the simulation.  It would also be worth investigating how sensitive some of our findings are w.r.t. the detailed chosen subgrid physics by probing alternative recipes and running higher  spatial resolution simulations.  The catalogues produced for this investigation (spectroscopical properties of the MareNostrum galaxies and its dark matter skeletons) are available online as discussed briefly in Appendix~\\ref{sec:catalogs}.  \\subsection*{Acknowledgements} { \\sl We thank F.~Brault for her help in  a preliminary investigation,  S.~Colombi, D.~Pogosyan, Y.~Dubois and D.~Aubert  for fruitful comments during the course of this work. This investigation was carried within the framework of the Horizon project, \\texttt{www.projet-horizon.fr}. The simulation was run on the MareNostrum machine at the Barcelona Supercomputing Centre and we would like to warmly thank the staff for their support and hospitality. We  also thank D.~Munro for freely distributing his Yorick programming language and opengl interface (available at {\\em\\tt http://yorick.sourceforge.net/}). CP thanks the  Leverhulme Trust for the visiting professorship F09846D. }"
    },
    "0910/0910.4427_arXiv.txt": {
        "abstract": "We report the results of our observations of H{\\sc i} absorption towards the central region of the rejuvenated radio galaxy 4C29.30 (J0840+2949) with the Giant Metrewave Radio Telescope (GMRT). The radio source has diffuse, extended emission with an angular size of $\\sim$520 arcsec (639 kpc) within which a compact edge-brightened double-lobed source with a size of 29 arcsec (36 kpc) is embedded. The absorption profile which is seen towards the central component of the inner double is well resolved and consists of six components; all but one of which appears to be red-shifted relative to the optical systemic velocity. The neutral hydrogen column density is estimated to be $N$(H{\\sc i})=4.7$\\times$10$^{21}$($T_s$/100)($f_c$/1.0) cm$^{-2}$, where $T_s$ and $f_c$ are the spin temperature and covering factor of the background source respectively. This detection reinforces a strong correlation between the occurrence of H{\\sc i} absorption and rejuvenation of radio activity suggested earlier, with the possibility that the red-shifted gas is fuelling the recent activity. ",
        "introduction": "One of the interesting and important questions in our understanding of active galactic nuclei (AGN) is whether such activity is usually episodic, and if so, the range of time scales of AGN activity. Besides helping to constrain models of episodic acitvity, this also has wider implications in our understanding of AGN feedback in structure formation and the evolution of galaxies (e.g. Sijacki et al. 2007, and references therein; Nesvadba \\& Lehnert 2008).  In the presently widely accepted paradigm, AGN activity is believed to be intimately related to the `feeding' of a supermassive black hole whose mass ranges from $\\sim$10$^6$ to 10$^{10}$ M$_\\odot$ (e.g. Marconi et al. 2004). Periodic `feeding' of the supermassive black hole may lead to different cycles of activity.  \\begin{figure*} \\vbox{ \\psfig{file=VLA1400collage.ps,width=6.35in,angle=00} } \\caption[]{A collage showing the large- and small-scale structure of the rejuvenated radio galaxy 4C29.30 (J0840+2949) at 1400 MHz reproduced from Jamrozy et al. (2007). The left panel shows the  D-array contour map of the entire source overlayed on the optical field from the Digital Sky Survey (DSS). The contour levels are spaced by factors of $\\sqrt{2}$ and the first contour is 0.3~mJy~beam$^{-1}$. The right panel shows the contour map of the central part of the source, containing the inner double, from the FIRST (Faint Images of the Radio Sky at Twenty-cm, Becker, White \\& Helfand 1995) survey. The contour levels are spaced by factors of $\\sqrt{2}$, and the first contour is 0.45~mJy~beam$^{-1}$. The size of the beam is indicated by an ellipse in the bottom right corner of each image. } \\end{figure*}  For radio-loud AGN, an interesting way of probing their history and hence episodic jet activity is via the structural and spectral information of the lobes of extended radio emission.  (e.g.  Burns, Schwendeman \\& White 1983; Burns, Feigelson \\& Schreier 1983; van Breugel \\& Fomalont 1984; Leahy, Pooley \\& Riley 1986; Baum et al. 1990; Clarke, Burns \\& Norman 1992; Junkes et al. 1993; Roettiger et al. 1994; Schoenmakers et al. 2000; Gizani \\& Leahy 2003; Konar et al. 2006). A very striking example of episodic jet activity is when a new pair of radio lobes is seen closer to the nucleus before the `old' and more distant radio lobes have faded (e.g. Subrahmanyan, Saripalli \\& Hunstead 1996; Lara et al. 1999). Such sources have been christened as `double-double' radio galaxies (DDRGs) by Schoenmakers et al. (2000). Saikia, Konar \\& Kulkarni (2006) reported the discovery of a new DDRG J0041+3224 and compiled a sample of approximately a dozen such objects from the literature, including 3C236 (Schilizzi et al. 2001) and J1247+6723 (Marecki et al. 2003; Bondi et al. 2004). The inner doubles in these two sources are compact with sizes of 1.7 kpc and 14 pc respectively, and have been classified as a compact steep spectrum (CSS) and a Gigahertz peaked spectrum (GPS) source respectively. The median linear size of the inner doubles for this sample of DDRGs is $\\sim$150 kpc, while the overall median total linear size is approximately a Mpc. In addition to the classic DDRGs, evidence of episodic activity may also be seen as diffuse, steep-spectrum emission from an earlier cycle of activity, in which a young double-lobed radio source may be embedded. An archetypal example of such a source, namely 4C29.30, was studied in detail by Jamrozy et al. (2007).  If the nuclear or jet activity is rejuvenated by a fresh supply of gas one might be able to find evidence of this gas via H{\\sc i} absorption towards the radio components in the central regions of the host galaxy. Saikia, Gupta \\& Konar (2007) reported the detection of H{\\sc i} absorption towards the inner double of the DDRG, J1247+6723. From the available information in the literature, they also suggested that there could be a strong relationship between the detection of H{\\sc i} gas and rejuvenation of radio activity.  We present the results of H{\\sc i} absorption towards the rejuvenated radio galaxy 4C29.30 with the Giant Metrewave Radio Telescope (GMRT). The total flux density of the inner double at 1287 MHz was estimated by Jamrozy et al. (2007) to be 390 mJy when observed with an angular resolution of $\\sim$2.6 arcsec, suggesting it to be a suitable source for these observations. We describe some of the basic properties of 4C29.30 (J0840+2949)  in Section 2, the observations and analyses in Section 3 and present the results and discussions in Section 4. The conclusions are summarised in Section 5. ",
        "conclusions": "Our GMRT image of the source (Fig. 2) with an angular resolution of 3.60$\\times$2.35 arcsec$^2$ along a position angle of 45.7$^\\circ$ and an rms noise of 0.4 mJy beam$^{-1}$ detects the inner double with a peak brightness of 78.9  mJy beam$^{-1}$ for the central component and 74.6 and 71.5 mJy beam$^{-1}$ for the northern and southern hotspots of the inner double. The peak brightness of the knot in the jet towards the south is 15.8 mJy beam$^{-1}$.  The H{\\sc i} absorption spectrum towards the central component is presented in Fig. 3.  H{\\sc i} absorption has been detected clearly towards the core of this rejuvenated radio galaxy. The rms noise in the spectrum is $\\sim$1 mJy beam$^{-1}$ channel$^{-1}$.  The absorption profile consists of a number of components all but one of which appear red-shifted relative to the optical systemic velocity.  The H{\\sc i} column density, $N$(H{\\sc i}), for the different spectral components has been estimated using the relation \\begin{equation} $N$({\\rm H{\\sc i}})=1.93\\times10^{18}\\frac{{T}_{s}~\\tau_p~\\Delta v}{f_c}~ {\\rm cm^{-2}}, \\label{eq1} \\end{equation} where $T_s$, $\\tau_p$, $\\Delta$$v$ and $f_c$ are the spin temperature, peak optical depth, the full width at half maximum (FWHM) of the Gaussian line profile, and the fraction of background emission covered by the absorber respectively. The estimates have been made assuming $T_s$=100 K and $f_c$=1.0.  The value of $T_s$ could be significantly greater than 100 K. For example, for the warm neutral medium seen in the Galaxy $T_s$ ranges from 5000$-$8000 K (Kulkarni \\& Heiles 1988). Such high spin temperatures are also expected to arise in the proximity of an active nucleus (Bahcall \\& Ekers 1969; Holt et al. 2006). The best fit to the spectrum with six Gaussian components is shown in Fig. 2 and the fit parameters are summarised in Table 1. The sixth component requires confirmation from more sensitive observations. The Table lists the optical heliocentric velocity, v$_{\\rm{hel}}$, FWHM of the Gaussian profile, the fractional absorption and the H{\\sc i} column density. The numbers within brackets are the errors on the quoted values. The  H{\\sc i} column density integrated over the entire absorption profile is 4.7$\\times$10$^{21}$ cm$^{-2}$.  The H{\\sc i} absorption spectra towards the hotspots and the `southern knot' are shown in Fig. 4.  No significant absorption has been detected towards the northern and sourthern hotspots of the inner double or the `southern knot' indicating a 3-$\\sigma$ upper limits to H{\\sc i} of $N$(H{\\sc i})=9.5$\\times$10$^{20}$ cm$^{-2}$, $N$(H{\\sc i})=8.3$\\times$10$^{20}$ cm$^{-2}$ and $N$(H{\\sc i})=5.5$\\times$10$^{21}$ cm$^{-2}$ respectively, assuming $T_s$=100 K, $f_c$=1.0 and $\\Delta$$v$=100 km s$^{-1}$. This indicates that the size of the absorber is less than $\\sim$35 kpc. However, the spectra of the `southern knot' and the northern hotspot show very marginal signs of absorption at $\\sim$1.948$\\times$10$^4$ km s$^{-1}$ which needs to be confirmed from more sensitive observations.  \\subsection{H{\\sc i} gas and rejuvenation of AGN activity} Saikia et al. (2007) explored any possible relationship between rejuvenation of radio or jet activity and the occurrence of H{\\sc i}. Unfortunately the number of sources is still small because most of the rejuvenated radio sources have weak radio emission in the central or nuclear region. In the list of DDRGs compiled by Saikia et al. (2006) the two exceptions are 3C236 and J1247+6723, with the flux density within a few kpc of the nuclear region being $\\gapp$100 mJy. The DDRG 3C236 which is a giant radio galaxy with a projected linear size of $\\sim$4250 kpc shows evidence of star formation and H{\\sc i} absorption against a lobe of the inner radio source (Conway \\& Schilizzi 2000; Schilizzi et al. 2001; O'Dea et al. 2001). Observations with milliarcsec resolution are required to determine the location of the absorbing clouds in the case of J1247+6723 reported by Saikia et al. (2007). An interesting case is 3C293, which could be classified as a misaligned DDRG, and exhibits absorption features both blue- and red-shifted relative to the the systemic velocity (Beswick, Pedlar \\& Holloway 2002; Beswick et al. 2004), and fast outflowing gas blue-shifted by upto $\\sim$1000 km s$^{-1}$ (Morganti et al. 2003; Emonts et al. 2005). The case for 3C258 (J1124+1919) as a rejuvenated galaxy is less clear. Although H{\\sc i} absorbing gas was reported by Gupta et al. (2006) towards the compact central source (e.g. Sanghera et al. 1995), we have not been able to confirm from GMRT observations (Saikia et al., in preparation) the weak extended emission reported by Strom et al. (1990). Another celebrated case is Centaurus A, where the inner double is due to more recent activity while the extended diffuse emission is due to earlier cycles of nuclear activity (Burns et al. 1983; Junkes et al. 1993; Ilana Feain, private communication), and H{\\sc i} absorption is seen towards the central region (e.g. Sarma, Troland \\& Rupen 2002; Morganti et al. 2008).  While the number of sources is still small, the detection of absorbing H{\\sc i} gas in the rejuvenated galaxies appears to be more frequent than for CSS and GPS objects (Vermeulen et al. 2003; Gupta et al. 2006), and the reported observations of 4C29.30 reinforces this trend. Although the sample size needs to be clearly increased, this trend is unlikely to be due to different source strengths. Considering the GPS objects listed by Gupta et al. (2006), which has the highest H{\\sc i} detection rate of $\\sim$45 per cent, the median total flux density of the sources at $\\sim$1400 MHz is $\\sim$2 Jy. Amongst the rejuvenated galaxies discussed here, the flux densities at $\\sim$1400 MHz of the entire central source of 3C236, 3C293 and Cen A are comparable to those of the GPS objects, while those of 4C29.30 and J1247+6723 are weaker than the GPS objects listed by Gupta et al. (2006). There should not be any bias because of source strength. Considering the H{\\sc i} column densities (Gupta et al. 2006), the median value for GPS sources is $\\sim$3$\\times$10$^{20}$ cm$^{-2}$. The rejuvenated radio galaxies discussed here have column densities in the range of $\\sim$8$-$50$\\times$10$^{20}$ cm$^{-2}$.   \\begin{table} \\caption{Multiple Gaussian fit to the H{\\sc i} absorption spectrum towards the central component (Fig. 3).} \\begin{center} \\begin{tabular}{|c|l|l|c|c|} \\hline Id. &   v$_{\\rm{hel}}$ &  FWHM      & Frac. abs. & $N$(H{\\sc i}) \\\\ no. &               &               & & 10$^{20}$  \\\\  % &   km s$^{-1}$ &   km s$^{-1}$ & & cm$^{-2}$ \\\\ \\hline 1  &  19383(2)  &   37(5)  &    0.072(0.008)   &    5.4(1.4)  \\\\ 2  &  19457(2)  &   33(3)  &    0.222(0.009)   &   16.1(2.0)  \\\\ 3  &  19490(1)  &   25(3)  &    0.243(0.013)   &   13.4(2.1)  \\\\ 4  &  19519(1)  &   18(2)  &    0.238(0.018)   &    9.3(1.6)  \\\\ 5  &  19549(4)  &   36(11) &    0.062(0.009)   &    4.4(2.0)  \\\\ 6  &  19599(1)  &   14(4)  &    0.065(0.013)   &    1.8(1.0)  \\\\ \\hline \\end{tabular} \\end{center} \\label{gauss} \\end{table}"
    },
    "0910/0910.1034_arXiv.txt": {
        "abstract": "{ Observations of dense molecular gas lie at the basis of our understanding of the density and temperature structure of protostellar envelopes and molecular outflows. The Atacama Pathfinder EXperiment (APEX) opens up the study of southern (Dec $<$ -35$^{\\circ}$) protostars. } {We aim to characterize the properties of the protostellar envelope, molecular outflow and surrounding cloud, through observations of high excitation molecular lines within a sample of 16 southern sources presumed to be embedded YSOs, including the most luminous Class I objects in Corona Australis and Chamaeleon.}  {Observations of submillimeter lines of CO, HCO$^+$ and their isotopologues, both single spectra and small maps (up to $80''\\times80''$), were taken with the FLASH and APEX-2a instruments mounted on APEX to trace the gas around the sources. The HARP-B instrument on the JCMT was used to map IRAS 15398-3359 in these lines. HCO$^+$ mapping probes the presence of dense centrally condensed gas, a characteristic of protostellar envelopes. The rare isotopologues C$^{18}$O and H$^{13}$CO$^+$ are also included to determine the optical depth, column density, and source velocity. The combination of multiple CO transitions, such as 3--2, 4--3 and 7--6, allows to constrain outflow properties, in particular the temperature. Archival submillimeter continuum data are used to determine envelope masses.} {Eleven  of the sixteen sources have associated warm and/or dense ($\\geq 10^6$ cm$^{-3}$) quiescent gas characteristic of protostellar envelopes, or an associated outflow.  Using the strength and degree of concentration of the HCO$^+$ 4--3 and CO 4--3 lines as a diagnostic,  five sources classified as Class I based on their spectral energy distributions are found not to be embedded YSOs. The C$^{18}$O 3--2 lines show that for none of the sources, foreground cloud layers are present.  Strong molecular outflows are found around six sources, with outflow forces an order of magnitude higher than for previously studied Class I sources of similar luminosity. }  {This study provides a starting point for future ALMA and Herschel surveys by identifying truly embedded southern YSOs and determining their larger scale envelope and outflow characteristics.} ",
        "introduction": "During the early stages of low-mass star formation, Young Stellar Objects (YSOs) are embedded in cold dark envelopes of gas and dust, which absorb the radiation from the central star \\citep{Lada87,Andre93}. This extinction is strong enough that low-mass embedded YSOs, or protostars, only emit weakly at infrared wavelengths \\citep[e.g.,][]{Jorgensen05,Gutermuth08}. Only at later evolutionary phases, in which the envelope has been accreted and/or dispersed, does emission in the optical and infrared (IR) dominate the Spectral Energy Distribution (SED) \\citep[e.g.][]{Hartmann05}. Protostars emit the bulk of their radiation at far-IR and sub-millimeter wavelengths, both as continuum radiation, produced by the cold ($T<30$ K) dust, and through molecular line emission from the gas-phase species present throughout the protostellar envelope. Although the bulk of the mass is accreted during the earliest embedded phases, more evolved protostellar envelopes still contain a reservoir of gas and dust that can accrete onto the central star and disk system and thus provide the material for disk and planet formation.  At the same time, jets and winds from the young star interact with the envelope and drive molecular outflows which clear the surroundings. Characterizing and quantifying all of these different physical components in the protostellar stage is still a major observational challenge.   The protostellar envelopes and molecular outflows can be directly observed either through thermal emission of dust at (sub)millimeter wavelengths \\citep[e.g.][]{Shirley00,Johnstone01,Nutter05} or the line emission of molecules. Low frequency molecular emission traces the cold gas in the protostellar envelopes \\citep[e.g.,][]{Hogerheijde98,Jorgensen02,Maret04} or molecular outflows \\citep[e.g.,][]{Snell90,Cabrit92,Bachiller99}. Single-dish observations of dust using current generation bolometer arrays are able to map large areas and image the surroundings of protostars \\citep[e.g.,][]{Motte98,Shirley00,Stanke06,Nutter08}. Through radiative transfer models \\citep[e.g.,][]{Shirley02,Jorgensen02,Young03}, including information from shorter wavelengths \\citep[e.g.,][]{Hatchell07}, the temperature and density structure of the protostellar envelope can be constrained, but the continuum data cannot determine the velocity structure of the infalling envelope, characterize the outflows and their interaction with the surroundings, or disentangle envelope and (foreground) cloud material. Analysis of gas observations in the form of spectra of multiple transitions of the same molecule and its isotopologues provide additional strong constraints on the physical characteristics of the protostellar envelope \\citep{Mangum93,Blake94,vanDishoeck95,Schoeier02,Maret04,Jorgensen05,Evans05,vanderTak07} and outflowing gas \\citep[e.g.,][]{Cabrit92,Bontemps96,Hogerheijde98,Hatchell99,Parise06,Hatchell07a}.   Although many different molecules have been observed in protostellar envelopes, only a limited number of species are well suited to trace the physical characteristics of all of the components of an embedded YSO and its surroundings.  For example, the use of CH$_3$OH and H$_2$CO is complicated by their changing abundances through the envelope, although some information can be obtained with careful analysis and a sufficient number of observed lines \\citep[e.g.][]{Mangum93, Blake94, vanDishoeck95, vanderTak00,Leurini04}. The weakness of high-excitation CH$_3$OH and H$_2$CO lines in all but a handful of the most luminous Class 0 protostars \\citep[e.g.][]{Jorgensen05b} coupled with a lack of some collisional rate coefficients make these species less suitable tracers for the bulk of the low-mass embedded YSOs. In practice, the column density of the surrounding cloud (and that of any unrelated foreground clouds) is best probed by low excitation optically thin transitions of molecules with a low dipole moment, such as the isotopologues of CO, e.g., C$^{18}$O 2--1 or 3--2 \\citep{vanKempen09}. For molecular outflows the line wings of various $^{12}$CO transitions have been efficiently used as tracers of their properties \\citep{Cabrit92,Bontemps96,Hogerheijde98,Bachiller99,Hatchell99,Hatchell07a}. The denser regions of the protostellar envelope need to be probed with emission lines with a high critical density, such as the higher excitation transitions from HCO$^+$ with critical densities $>10^6$ cm$^{-3}$. The warm gas ($T>$50 K), which can be present in both the molecular outflow and the inner region of the protostellar envelope, can only be traced by spectrally resolved high-$J$ CO transitions, which have $E_{\\rm{up}}>50 $ K for $J_{\\rm{up}}\\geq 4$ \\citep{Hogerheijde98}.   Embedded YSOs are generally identified by their 2-24 $\\mu$m IR slope with positive values characteristic of Class I objects \\citep{Lada87}. Over the last several years, it has been found that some of these Class I objects turn out to be edge-on disks or obscured sources \\citep[e.g.,][]{Luhman99,Brandner00,Pontoppidan05,Lahuis06,vanKempen09}. The use of a spectral line map over a small ($\\sim$2$'\\times2'$) region around the protostar is an elegant solution to disentangle the contributions of the different components \\citep{Boogert02}.  In particular, the spatial distribution of the HCO$^+$ 4--3 line has proven to be an excellent diagnostic of such `false' embedded sources: in the case of Ophiuchus, about 60\\% of the sources with Class I or flat  SED slopes turned out not to be embedded YSOs \\citep{vanKempen09}.   With the development of the Atacama Large Millimeter/Submillimeter Array (ALMA) on Chajnantor, Chile, surveys of sources in southern star-forming clouds will be undertaken at high resolution ($\\leq$ 1$''$) at a wide range of frequencies, and are expected to reveal much about the inner structure of protostars and their circumstellar disks. The {\\it Herschel Space Observatory} will also target a large number of low-mass YSOs across the sky.  Both sets of observations rely heavily on complementary large aperture ground-based single-dish observations for planning and interpretation.  Apart from providing the necessary information about the structure on larger scales, single-dish studies of low-mass star formation in the southern sky will also put the results obtained from studies on the northern hemisphere in perspective.   Due to a lack of sub-millimeter telescopes at high dry sites in the southern hemisphere, high frequency data on southern YSOs are still limited. The Swedish ESO Submillimeter Telescope (SEST) operated up to the 345 GHz window, but for only limited periods of the year.  As a result, southern clouds, such as Chamaeleon, Corona Autralis, Vela or Lupus, are much less studied than their counterparts in the northern sky, such as Taurus and Perseus, where such observations have been readily available for over a decade \\citep[e.g.,][]{Hogerheijde97,Motte98,Hogerheijde98,Johnstone00,Jorgensen02,Nutter05,Nutter08}. The Atacama Pathfinder EXperiment (APEX) \\citep{Guesten06} \\footnote{This publication is based on data acquired with the Atacama Pathfinder Experiment (APEX) with programs E-77.C-0217, E-77.C-4010 and E-78.C-0576. APEX is a collaboration between the Max-Planck-Institut f\\\"ur Radioastronomie, the European Southern Observatory, and the Onsala Space Observatory.}  has opened up access to the atmospheric windows in the 200-1400 $\\mu$m wavelength regime over the entire southern sky.  Of the southern star-forming regions not or only poorly visible with northern telescopes two regions have proven especially interesting. The Chamaeleon I cloud ($D$=130 pc, Dec = -77$^\\circ$), observed with {\\it IRAS} and the {\\it Infrared Space Observatory} (ISO) and included in the {\\it Spitzer Space Telescope} guaranteed time and $`$Cores to disks' (c2d) Legacy programs \\citep{Persi00,Evans03,Damjanov07,Luhman08} contains some embedded sources, especially around the Cederblad region \\citep{Persi00,Lehtinen01,Belloche06,Hiramatsu07}.  Corona Australis is a nearby star-forming region ($D$=170 pc, see \\citet{Knude98}; Dec = -36$^\\circ$), well-known for the central Coronet cluster near the R CrA star \\citep{Loren79,Taylor84,Wilking86,Brown87}. Large-scale C$^{18}$O 1--0 maps show that it contains about 50 M$_\\odot$ of gas and dust \\citep{Harju93}. Many surveys have been undertaken in the IR \\citep[e.g.,][]{Wilking86, Wilking97,Olofsson99,Nisini05}, but only a few studies mapped this region at submillimeter wavelengths \\citep{Chini03,Nutter05}. Although Corona Australis has only a few protostars with rising infrared SEDs, the cloud does contains some of the most luminous low-mass protostars in the neighborhood of the Sun ($D<$ 200 pc) with luminosities of up to 20 L$_{\\rm{\\odot}}$.   In this paper we present observations of submillimeter lines of CO and HCO$^+$ of a sample of 16 embedded sources in the southern sky, with a focus on the Chamaeleon I and Corona Australis clouds, to identify basic parameters such as column density, presence and influence of outflowing material, presence of warm and dense gas and the influence of the immediate surroundings, in preparation for Herschel and ALMA surveys or in-depth high resolution interferometric observations.  In section $\\S$ 2 we present the observations, for which the results are given in $\\S$ 3 with a distinction between the clouds and isolated sources. Both the single spectra ($\\S$ 3.1) and maps ($\\S$ 3.2) are presented.  In $\\S$ 4 we perform the analysis of the observations, making use of archival submillimeter continuum data, with the final conclusions given in $\\S$ 5.  \\placeTableChapterFourOne ",
        "conclusions": "We present observations of CO, HCO$^+$ and their isotopologues, ranging in excitation from CO 3--2 to CO 7--6 of a sample of southern embedded sources to probe the molecular gas content in both the protostellar envelopes and molecular outflows in preparation for future ALMA and Herschel surveys. The main conclusions are the following: \\begin{itemize} \\item HCO$^+$ 4--3 and CO 4--3 integrated intensities and concentrations, combined with information on the presence of outflows, confirm the presence of warm dense quiescent gas associated with 11 of our 16 sources.  \\item RCrA TS 3.5, Cha IRS 6a, Cha INa2, IRAS 07178-4429 and IRAS 13546-3941 are likely not embedded YSOs due to the lack of HCO$^+$ 4--3 emission and/or the lack of central concentration in HCO$^+$ 4--3.   \\item The swept-up outflow gas has temperatures of the order of 50--100 K, as measured from the ratios of the CO line profile wings, with the values depending on the adopted ambient densities. These values are comparable to the temperatures predicted by the heating model of \\citet{Hatchell99}, with the exception of IRAS 15398-3359, which may be unusually warm.  \\item The outflows of 6 of our truly embedded sources --- Ced 110 IRS 4, CrA IRAS 32, RCrA IRS 7A, HH 100, HH 46,IRAS 15398-3359 --- were characterized using CO spectral maps. The outflows all have exceptionally strong outflow forces, almost two orders of magnitude higher than expected from their luminosities following the relation of \\citet{Bontemps96}.  \\item Neither Chamaeleon nor Corona Australis have foreground layers as found in Ophiuchus L~1688. All C$^{18}$O 3--2 spectra can be fitted with single gaussians.   \\end{itemize}  Future observations using ALMA and Herschel will be able to probe the molecular emission associated with these YSOs at higher spatial resolution and higher frequencies. Comparison with single dish data, both continuum surveys and spectral line mapping, will be essential in analyzing the results obtained with future facilities.\\\\"
    },
    "0910/0910.5220_arXiv.txt": {
        "abstract": "We generalize the ``renormalized'' perturbation theory (RPT) formalism of \\citet{CrocceScoccimarro2006a} to deal with multiple fluids in the Universe and here we present the complete calculations up to the one-loop level in the RPT. We apply this approach to the problem of following the non-linear evolution of baryon and cold dark matter (CDM) perturbations, evolving from the distinct sets of initial conditions, from the high redshift post-recombination Universe right through to the present day.  In current theoretical and numerical models of structure formation, it is standard practice to treat baryons and CDM as an effective single matter fluid -- the so called dark matter only modeling. In this approximation, one uses a weighed sum of late time baryon and CDM transfer functions to set initial mass fluctuations. In this paper we explore whether this approach can be employed for high precision modeling of structure formation. We show that, even if we only follow the linear evolution, there is a large-scale scale-dependent bias between baryons and CDM for the currently favored WMAP5 $\\Lambda$CDM model. This time evolving bias is significant $(>1\\%)$ until the present day, when it is driven towards unity through gravitational relaxation processes. Using the RPT formalism we test this approximation in the non-linear regime. We show that the non-linear CDM power spectrum in the 2-component fluid differs from that obtained from an effective mean-mass 1-component fluid by $\\sim3\\%$ on scales of order $k\\sim0.05\\kMpc$ at $z=10$, and by $\\sim0.5\\%$ at $z=0$. However, for the case of the non-linear evolution of the baryons the situation is worse and we find that the power spectrum is suppressed, relative to the total matter, by $\\sim15\\%$ on scales $k\\sim0.05\\kMpc$ at $z=10$, and by $\\sim3-5\\%$ at $z=0$. Importantly, besides the suppression of the spectrum, the Baryonic Acoustic Oscillation (BAO) features are amplified for baryon and slightly damped for CDM spectra. If we compare the total matter power spectra in the 2- and 1-component fluid approaches, then we find excellent agreement, with deviations being $<0.5\\%$ throughout the evolution.  Consequences: high precision modeling of the large-scale distribution of baryons in the Universe can not be achieved through an effective mean-mass 1-component fluid approximation; detection significance of BAO will be amplified in probes that study baryonic matter, relative to probes that study the CDM or total mass only. The CDM distribution can be modeled accurately at late times and the total matter at all times. This is good news for probes that are sensitive to the total mass, such as gravitational weak lensing as existing modeling techniques are good enough. Lastly, we identify an analytic approximation that greatly simplifies the evaluation of the full PT expressions, and it is better than $<1\\%$ over the full range of scales and times considered. ",
        "introduction": "\\label{sec:intro}  In the current paradigm for structure formation in the Universe, it is supposed that there was an initial period of inflation, during which, quantum fluctuations were generated and inflated up to super-horizon scales; producing a near scale-invariant and near Gaussian set of primordial potential fluctuations. At the end of inflation the Universe is reheated, and particles and radiation are synthesized. In this hot early phase, cold thermal relics are also produced, these particles interact gravitationally and possibly through the weak interaction -- dubbed Cold Dark Matter (CDM).  Prior to recombination photons are coupled to electrons through Thompson scattering, and in turn electrons to baryonic nuclei through the Coulomb interaction. Baryon fluctuations are pressure supported on scales smaller than the sound horizon scale.  Subsequent to recombination photons free-stream out of perturbations and baryons cool into the CDM potential wells. Structure formation then proceeds in a hierarchical way with small objects collapsing first and then merging to form larger objects.  Eventually, sufficiently dense gas wells are accumulated and galaxies form.  At late times the Universe switches from a decelerating phase of expansion to an accelerated phase. This is attributed to the non-zero constant energy-density of the vacuum. This paradigm has been dubbed: $\\Lambda$CDM \\citep[][]{PeeblesRatra2003,Spergeletal2007,Komatsuetal2009}.  The present day energy-density budget for the $\\Lambda$CDM model is distributed into several components, which in units of the critical density are: vacuum energy $\\Omega_{\\Lambda,0}\\approx0.73$, matter $\\Omega_{\\rmm,0}\\approx0.27$, neutrinos $\\Omega_{\\nu,0}\\lesssim 10^{-2}$ and radiation $\\Omega_{\\rr,0}\\approx 5\\times 10^{-4}$. The matter distribution can be subdivided further into contributions from CDM and baryons, with present day values $\\Omega_{\\rc}\\approx0.225$ and $\\Omega_{\\rb}\\approx0.045$. The detailed physics for the evolution of the radiation, CDM, baryon and neutrino fluctuations ($\\delta^{\\rr}, \\delta^{\\rc}, \\delta^{\\rb}, \\delta^{\\nu}$), from the early Universe through to recombination can be obtained by solving the Einstein--Boltzmann equations. These are a set of coupled non-linear partial differential equations, however while the fluctuations are small they may be linearized and solved \\citep[for a review see][]{MaBertschinger1995,SeljakZaldarriaga1996}.  At later times these fluctuations enter a phase of non-linear growth. Their evolution during this period can no longer be described accurately using the linear Einstein--Boltzmann theory and must be followed using higher order perturbation theory (hereafter PT) techniques \\cite{Bernardeauetal2002} or more directly through $N$-body simulations \\cite{Bertschinger2001}.   \\begin{figure}[t!] \\centering{ \\includegraphics[width=8cm,clip=]{Fig.1.eps}} \\caption{\\small{Baryon bias as a function of inverse spatial scale, where we have defined the baryon bias $b_{\\rb}(k,z)\\equiv \\delta^{\\rb}_{\\mathrm{lin}}(k,z)/\\delta^{\\rc}_{\\mathrm{lin}}(k,z)$, and used \\Eqns{eq:Lin_c}{eq:Lin_b} from \\S\\ref{ssec:application}. Solid through to dashed lines show results for redshifts $z=\\{0, 1.0, 3.0, 5.0, 10.0, 20.0\\}$.}\\label{fig:baryonbias}} \\end{figure}   The cross-over from the linear to the non-linear theory is, in general, not straightforward.  Most of the theoretical and numerical approaches to this latter task were developed during the previous two decades, during which time the Standard CDM model was favored: essentially Einstein-de Sitter spacetime, $\\Omega_{\\rmm}=1$, with energy-density dominated by CDM $\\gtrsim 97\\%$, and baryons contributing $\\lesssim3\\%$ to the total matter. For this case it is natural to assume that the CDM fluctuations dominate the gravitational potentials at nearly all times, with baryons playing little role in the formation of large-scale structure \\citep[see for example][]{Efstathiouetal1985,ThomasCouchman1992,Couchmanetal1995}. Thus, one simply requires a transfer function for the CDM distribution at some late time, to model the evolution of both components.  In the latter part of the 1990s, cosmological tests pointed towards the $\\Lambda$CDM paradigm, and it was recognized that baryons should play some role in shaping the transfer function of matter fluctuations \\citep{MaBertschinger1995,Sugiyama1995,EisensteinHu1998,MeiksinWhitePeacock1999}. However, rather than attempting to follow the CDM and baryons as separate fluids evolving from distinct sets of initial conditions, a simpler approximate scheme was adopted. This scheme is currently standard practice for all studies of structure formation in the Universe \\citep{Pearceetal2001,Teyssier2002,Springel2005,Springel2009}. It can be summarized as follows: \\begin{enumerate} \\item Fix the cosmological model, specifying $\\Omega_{\\rc}$ and $\\Omega_{\\rb}$, and hence the fraction of baryons, $\\fb$. Solve for the evolution of all perturbed species using the linearized coupled Einstein--Boltzmann equations. Obtain transfer functions at $z=0$, where CDM and baryons are fully relaxed: i.e.~$T^{\\rc}(k,z=0)\\approx T^{\\rb}(k,z=0)$, where $T^{\\rc}$ and $T^{\\rb}$ are the transfer functions of the CDM and baryons, respectively. \\item Use the transfer functions to generate the linear matter power spectrum, normalized to the present day: i.e. $P_{\\bar{\\delta}\\bar{\\delta}}(\\bk,z=0)\\approx \\left[(1-\\fb) T^{\\rc}(k,z=0)+\\fb T^{\\rb}(k,z=0)\\right]^2Ak^n \\approx [T^{\\rc}(k,z=0)]^2 Ak^n\\ ,$ where the total matter fluctuation is $\\bar{\\delta}=(1-\\fb)\\delta^{\\rc}+\\fb\\delta^{\\rb}$. \\item Scale this power spectrum back to the initial time $z_i$, using the linear growth factor for the total matter fluctuation $\\bar{\\delta}$, which obtains from solving the equations of motion for a single fluid. \\item Generate the initial CDM density field and assume that the baryons are perfect tracers of the CDM. \\item Evolve this effective CDM+baryon fluid under gravity using full non-linear equations of motion, either as a single fluid for dark matter only, or as a 2-component fluid if hydrodynamics are also to be followed. \\end{enumerate}  This effective description for the formation of structure becomes poor as the baryon fraction $\\fb=\\Omega_{\\rb}/\\Omega_{\\rmm}$ approaches $\\sim O(1)$, and for the currently favored WMAP5 cosmology $\\fb\\approx0.17$. In fact, in linear theory the gravitational relaxation between baryonic and dark matter components is only achieved at the level of $<1\\%$ by $z=0$.  Moreover, as can be seen in \\fig{fig:baryonbias} there is a non-trivial scale-dependent bias between baryons and dark matter that exists right through to the present day \\footnote{Note that we have defined the baryon bias $b_{\\rb}(k,z)\\equiv\\delta^{\\rb}_{\\mathrm{lin}}(k,z)/ \\delta^{\\rc}_{\\mathrm{lin}}(k,z)$, where these linear fluctuations are given by \\Eqns{eq:Lin_c}{eq:Lin_b} from \\S\\ref{ssec:application}. In this analysis we have neglected any residual effect of the coupling between the photons and the baryons, and have assumed that the baryons are cold at $z=100$, hence $p_{\\rm b}\\propto \\rho T\\sim 0$. Furthermore, this result is sensitive to our choice for the initial $\\delta_i$ and $\\theta_i$, and indeed had we chosen the initial baryon and CDM velocity fields to be proportional to the {\\em total matter} fluctuation, then these effects would be slightly reduced. However, the resultant bias will still remain non-negligible.}. Failing, to take these biases into account may lead to non-negligible systematic errors in studies that attempt to obtain cosmological constraints from Large-Scale Structure (LSS) tests. With the next generation of LSS tests aiming to achieving 1\\% precision in measurements of the matter power spectrum \\cite{DETF2006,ESO2006}, it seems timely to investigate whether such modeling assumptions may impact inferences. Where we expect such systematic effects to play an important role is for studies based on baryonic physics such as: using 21cm emission from neutral hydrogen to trace mass density at the epoch of reionization $z_{\\rm reion} \\approx 10$ \\citep{Zhangetal2004,Ilievetal2006,Furlanettoetal2006,PillepichPorciani2007}; and studies of the Lyman alpha forest to probe the matter power spectrum \\citep{Croftetal1998,McDonaldetal2006}.  In this, the first in a series of papers, we test whether this effective procedure can indeed be employed to follow, at high precision, the non-linear growth of structure formation in baryon and CDM fluids, from recombination right through to the present day. We do this by generalizing the matrix based ``renormalized'' perturbation theory (RPT) \\footnote{It appears to us that what may be implemented in the RPT framework is not so much what would usually be understood as the ``perturbative renormalization'' of the theory, but rather a ``resummation'' of the perturbative series, in some specific approximation (the ``high-$k$'' limit). We will continue to use the acronym RPT to refer to the method, but would suggest that it is more properly read as {\\em resummed}, and not {\\em renormalized} perturbation theory.} techniques of \\citet{CrocceScoccimarro2006a,CrocceScoccimarro2006b,CrocceScoccimarro2008} to an arbitrary number of gravitating fluids. In this paper we content our selves with calculating the observables up to the one-loop level in the theory. In a future paper we consider the resummation to arbitrary orders (Somogyi \\& Smith 2009, in preparation).  The late-time linear theory evolution of coupled baryon and CDM fluctuations, with sudden late time reionization of the inter-galactic medium was studied by \\citep{Nusser2000,MatarreseMohayaee2002}. Recently, \\citep{ShojiKomatsu2009} explored a similar system using non-linear PT techniques. The work presented by these authors differs from that presented here in a number of fundamental ways. Firstly, these authors have assumed that, on large scales, baryons and CDM fluctuations evolve with the same initial conditions and that the only physical difference between the two fluids arises on small scales, where galaxy formation processes are able to provide some pressure support. They model this with a fixed comoving-scale Jeans criterion. In this work we are not so much interested in following how small-scale baryonic collapse feeds back on the mass distribution, but are more interested in following the very large-scale baryon and CDM fluctuations, initialized with their correct post-recombination spatial distributions, into the non-linear regime. Secondly, their work was developed within standard PT, the problems associated with this formalism are well documented \\citep{Bernardeauetal2002}. Here, we develop the more successful RPT approach \\citep{CrocceScoccimarro2006a}.  In passing, we note that a number of additional complimentary approaches and extensions to single-fluid RPT have recently been developed \\citep{MatarresePietroni2007,MatarresePietroni2008,Bernardeauetal2008,Pietroni2008,Matsubara2008}. In addition, there has recently been progress on realistic modeling of massive neutrinos and CDM \\citep{Wong2008,Saitoetal2008,Ichikietal2009,Brandbygeetal2008,BrandbygeHannestad2009a,BrandbygeHannestad2009b,Lesgourguesetal2009}, for simplicity we shall assume that neutrinos play a negligible role.  The paper is broken down as follows: In \\Sect{sec:EoM} we set up the theoretical formalism: presenting the coupled equations of motion for the $N$-component fluid. Then we show how, when written in matrix form, the equations may be solved at linear order. We discuss generalized initial conditions, and then discuss the specific case of baryons and CDM. In \\Sect{sec:PT} we develop the non-linear PT expansion for the solutions, and show that they may be constructed to arbitrary order using a set of Feynman rules.  In \\Sect{sec:stats} we discuss the two-point statistics of the fields, which we take as our lowest order observables.  In \\Sect{sec:OneLoopPower} we give specific details of how we calculate the one-loop corrections in our theory. In \\Sect{sec:results} we present results for various quantities of interest. In \\Sect{sec:approx} we give details of an approximate RPT scheme, which matches the exact calculation to within $<1\\%$, for all times of interest and on all scales where the perturbation theory is valid. Finally, in \\Sect{sec:conclusions} we draw our conclusions. ",
        "conclusions": "\\label{sec:conclusions}  In this paper we have generalized the matrix based RPT formalism of \\citet{CrocceScoccimarro2006a} to the problem of dealing with $N$-component fluids, each of which evolves from a distinct set of initial conditions and contributes to the matter-density and velocity fields.  In \\Sect{sec:EoM} we showed that, for $N$-component fluids, the essential structure of the RPT framework remained the same as for the $N=1$ case -- that is, the equations of motion could be solved in exactly the same way. However, the details of the building blocks of the theory did change: i.e.~the vertex, propagator and initial conditions. For $N=2$ the vertex matrix offered no surprises. The linear propagator was more interesting, and we showed that besides the usual growing and decaying modes $\\{D_{+}\\propto \\re^{\\eta}, D^{(1)}_{-}\\propto \\re^{-3\\eta/2}\\}$, there were two additional eigenvalues: a static and decaying mode $\\{D_{\\rm stat}\\propto {\\rm const}, D^{(2)}_{-}\\propto \\re^{-\\eta/2}\\}$. Interestingly, in going to $N>2$ no further time dependence was gained. Unlike the 1-component fluid case, it was no longer possible to choose pure growing modes, unless all of the fluids evolved from identical initial conditions. However, we showed that one can still set up the initial conditions to possess pure large-scale growing modes. All our results were obtained for this specific choice of initial conditions.  In \\Sect{sec:PT} it was shown that the equations of motion for the $N$-component fluid case could be solved using a power series solution and that the solutions could be recovered order by order. A graphical representation of the solutions was also described.  In \\Sect{sec:stats} we discussed the statistics of the fields: in particular the covariance matrix of the $N$-component fluid perturbations, and the fully non-linear propagator.  This latter quantity can be understood to be the statistic that provides information on the memory of the evolved field to the initial conditions. It also possesses a perturbative expansion and we gave expressions up to one-loop order. A diagrammatic representation was also discussed.  As for the case of the 1-component fluid RPT, we then showed that for the $N$-component fluid RPT the full non-linear power spectrum also possessed a series expansion, and that the series could be grouped into two sub-series, ``reducible'' and ``mode-coupling''. It was shown that the reducible terms were proportional to the initial power spectrum and non-linear propagators. Again we gave complete details for the series up to the one-loop level in the expansions. A diagrammatic description was also discussed.  In \\Sect{sec:OneLoopPower} we developed further the one-loop expressions for the propagator and the mode-coupling contribution to the power spectrum. These expressions were then given for the explicit case of a 2-component fluid.  In \\Sect{sec:results} we applied the formalism to the problem of modeling the non-linear evolution of a CDM and baryon fluid in the $\\Lambda$CDM paradigm. This enabled us to test the validity of the approximation of treating CDM and baryons as an effective 1-component fluid evolving from a single set of initial conditions, as is currently standard practice.  For the case of CDM, it was shown that the approximation was very good for $z<3$, with the exact and approximate CDM power spectra differing by $<1\\%$. However, for higher redshifts the approximation became progressively worse, it being of the order $\\sim3\\%$ at $z=10$, with the power in the exact theory being amplified and having weaker BAO features than the approximate theory.  For the case of baryons the situation was worse: at $z=20$ the exact and approximate spectra differed by $\\sim25\\%$, and at $z=10$ by $\\sim15\\%$. At later times the approximation was still quite poor, and at $z=3$ the errors were of the order $\\sim 5\\%$, and by $z=0$ were still between $\\sim2-3\\%$. The power in the exact theory was suppressed in amplitude but possessed stronger BAO than the corresponding power for the approximate theory.  These conclusions remained essentially unaffected, when the so-called approximate 1-component fluid initial conditions were employed for the effective 1-component fluid computation.  Lastly, we computed the total non-linear matter power spectrum and found that the exact and approximate theory agreed to within $<0.03\\%$, when the so-called exact 1-component fluid initial conditions were used. Employing the approximate 1-component fluid initial conditions on the other hand leads to an agreement to within $<0.5\\%$ on scales where the perturbation theory was valid $k<0.2\\kMpc$ and for all times of interest. Deviations on the order of $\\sim0.7\\%$ were found on smaller scales with this approximation.  \\vspace{0.2cm}  Our main conclusions therefore are:  \\begin{itemize} \\item For theoretical modeling of the low-redshift CDM distribution, the approximate 1-component fluid treatment provides a good approximation. For higher redshift studies, such as those that probe the epoch of reionization, or attempt to model the first objects that form (stars/haloes), then one must be careful to use the appropriate CDM transfer function for the redshift of interest. Otherwise, systematic errors will be present at the level of several percent. \\item For theoretical modeling of baryons in the Universe, a 1-component fluid treatment leads to significant $>1\\%$ errors at all times. This implies that cosmological probes that are primarily sensitive to the distributions of baryons in the universe, such as: 21 cm radiation from neutral hydrogen, or the Lyman alpha forest absorption lines in high redshift quasars, can not be modeled using dark matter only simulations at high precision. Instead the baryons must be modeled using a 2-component fluid system of CDM and baryons, with the initial distributions being specified by the linear Einstein--Boltzmann solutions. The correct treatment of CDM and baryon initial conditions may also affect how galaxies form. \\item Observational probes for cosmology that are sensitive to the total mass distribution, such as weak gravitational lensing, can be accurately interpreted using an effective single CDM plus baryon fluid. Thus current modeling technology is good enough for high precision work, but it is highly desirable to use the exact prescription when specifying the initial conditions for the effective single fluid. \\end{itemize}"
    },
    "0910/0910.5685_arXiv.txt": {
        "abstract": "When scientific experiments require transmission of powerful laser or radio beams through the atmosphere the Federal Aviation Administration (FAA) requires that precautions be taken to avoid inadvertent illumination of aircraft. Here we describe a highly reliable system for detecting aircraft entering the vicinity of a laser beam by making use of the Air Traffic Control (ATC) transponders required on most aircraft. This system uses two antennas, both aligned with the laser beam. One antenna has a broad beam and the other has a narrow beam. The ratio of the transponder power received in the narrow beam to that received in the broad beam gives a measure of the angular distance of the aircraft from the axis that is independent of the range or the transmitter power. This ratio is easily measured and can be used to shutter the laser when the aircraft is too close to the beam.  Prototype systems operating on astronomical telescopes have produced good results. ",
        "introduction": "A number of scientific experiments require the transmission of a laser beam through the atmosphere, using an astronomical telescope or its equivalent. In order to avoid hazard to aircraft the FAA requires that one or more observers be stationed outside any telescope that is transmitting a laser beam. These observers close the laser shutter when an aircraft is observed within $25^\\circ$ of the laser beam (as viewed from the telescope). Such experiments include: lunar and satellite laser ranging \\citep{llr,ilrs}; creation of artificial guide stars for active optics \\citep[e.g.,][]{keck-lgs}; and atmospheric remote sensing using lidar \\citep{lidar1,lidar2}. In this paper we discuss part of an aircraft detection system now employed at the Apache Point Observatory (APO) for a lunar ranging experiment called APOLLO \\citep{apollo}. This system could be used not only for the other laser beam experiments mentioned above, but also for protecting aircraft from high-powered radar transmitters such as those used for ionospheric research \\citep[e.g.,][]{ionosphere}.  The detection scheme described here is used in conjunction with a complementary infrared camera detection system, together providing robust protection to aircraft.  The FAA rules effectively require transponders on all commercial aircraft and most private aircraft \\citep[the exact language can be found in Section 91.215 of the Federal Aviation Regulations:][]{far}.  These transponders are interrogated frequently (at 1030~MHz) by the regional ATC radars and also by the airborne Traffic Collision Avoidance System (TCAS). The transponders reply incoherently at $1090\\pm 3$~MHz with a pulse coded response. The response must have vertical electric field polarization, an omni-directional pattern, and transmitted peak power between 70 and 500~W. Various coding schemes convey information about the aircraft.  Mode-A and Mode-C responses communicate a temporarily assigned aircraft identity and altitude, respectively.  A newer Mode-S encoding flexibly communicates permanent aircraft identity, coordinates, altitude, etc., but still as pulsed transmission at 1090~MHz.  The APOLLO laser is never used at elevations less than $15^\\circ$ and this elevation restriction is typical of other laser and radar transmitters. Thus for practical altitudes of $< 13$~km the aircraft range will not exceed $\\sim$50~km, and the received power is very high by modern communications standards ($> -69$~dBm) so that it may be easily detected with a total power receiver. However the received power is highly variable because both the range and the transmitted power are variable.  The design requirement is a highly reliable method of detecting when an aircraft transponder is within about $15^\\circ$ of the telescope beam. This $15^\\circ$ specification---differing from the 25$^\\circ$ angle used by human spotters---is set by the expected angular rate of aircraft, transponder interrogation frequency, and the desire to avoid excessive triggers when pointing the beam as low as $15^\\circ$ above the horizon. The general concept is to use two antennas aligned with the optical axis of the telescope, one with a beam width (full-width at half-power) of about $30^\\circ$ and the other with a beam width of about $90^\\circ$, as shown in Figure~\\ref{fig:gains}. The ratio of the power received by the narrow beam antenna to that received by the broad beam antenna depends only on the angular position of the transponder with respect to the beam axis. In particular it does not depend on the distance, transmitted power, or polarization mismatch.  \\begin{figure} \\begin{center}\\includegraphics[width=89mm, angle=0]{fig1.pdf} \\end{center}  \\caption{Gain of the broad-beam and narrow-beam antennas as a function of the angular separation of the transponder from the axis. The solid lines are H-plane cuts through the gains of the antennas actually used and the dashed lines are E-plane cuts. If the antennas are pointed to zero elevation, the E plane is the vertical plane and the H plane is horizontal.  The power received is proportional to the gain. \\label{fig:gains}}  \\end{figure}  The APO telescope, like many others used in this type of experiment, is on an altazimuth mount with a secondary mirror more than 80~cm in diameter centered on the optical axis.  The radio antennas can be mounted facing the sky on the secondary support structure and aligned with the optical axis without interfering with the optical beam, and the polarization will remain vertical as the telescope is moved. This geometry is also typical of the radar antennas used for ionospheric research although the secondary reflector is much larger in these cases.  For telescopes on an equatorial mount the position angle of the linear polarization changes with hour angle. This variation can be easily accommodated by changing to antennas that are sensitive to circular polarization.  Simple patch antennas are well-suited to this application because the narrow bandwidth of a simple patch is an advantage when the signal also has a narrow bandwidth.  Patches are also very robust mechanically---a significant advantage in this application.  The polarization of a patch can be changed from linear to circular simply by moving the feed point.  A single patch is suitable for the broad beam antenna and the narrow beam antenna can be made with an array of patches. The ratio of the power in the array to the power in a single patch will depend only on the array factor, which is easily calculated. The element spacing of the array can be adjusted to optimize the beam width and the sidelobe levels.  The system design consists primarily of: impedance matching a simple patch antenna at 1090~MHz; adjusting the array configuration to obtain a suitable beam width and adequately low sidelobes; designing total power detectors for the two channels; development of signal processing electronics; and devising a reliable calibration system.  A block diagram of the analog signal flow is shown in Figure~\\ref{fig:block_diagram}. Here one can see that we have used the center element of the array both as an array element and as the broad beam element by splitting the signal with a power divider. To compensate for that power division, the other array elements must have -3~dB attenuators before the array summer.  \\begin{figure} \\begin{center}\\includegraphics[width=170mm, angle=0]{fig2.pdf} \\end{center}  \\caption{Block diagram of the analog signal flow. The patches are shown with the E field vertical.  The azimuthal angle is defined with respect to the horizontal. The array is drawn approximately to scale. The center patch is used both as an array element and as the broad beam element. \\label{fig:block_diagram}}  \\end{figure} ",
        "conclusions": ""
    },
    "0910/0910.2546_arXiv.txt": {
        "abstract": "Most stars are members of binaries, and the evolution of a star in a close binary system differs from that of an ioslated star due to the proximity of its companion star. The components in a binary system interact in many ways and binary evolution leads to the formation of many peculiar stars, including blue stragglers and hot subdwarfs. We will discuss binary evolution and the formation of blue stragglers and hot subdwarfs, and show that those hot objects are important in the study of evolutionary population synthesis (EPS), and conclude that binary interactions should be included in the study of EPS. Indeed, binary interactions make a stellar population younger (hotter), and the far-ultraviolet (UV) excess in elliptical galaxies is shown to be most likely resulted from binary interactions. This has major implications for understanding the evolution of the far-UV excess and elliptical galaxies in general. In particular, it implies that the far-UV excess is not a sign of age, as had been postulated prviously and predicts that it should not be strongly dependent on the metallicity of the population, but exists universally from dwarf ellipticals to giant ellipticals. ",
        "introduction": "Evolutionary population synthesis (EPS) has experienced a rapid progress since the early 90's and provides the most robust approach in studying stellar populations of galaxies. In most of the current EPS models, binary evolution has been ignored. However, most stars are members of binaries, and binary evolution leads to the formation of many peculiar objects, such as blue stagglers and hot subdwarfs, which are hot, long-lived and still very luminous. Those objects contribute very much to the spectral energy distribution (SED) at short wavelength for an old stellar population, and make the population look hotter and younger. Such a ``cosmetic'' effect has been successfully included by EPS models of \\cite[Zhang et al.\\ (2004)]{zha04}, \\cite[Han, Podsiadlowski \\& Lynas-Gray (2007)]{han07}, \\cite[Chen \\& Han (2009)]{che09}. Some puzzles in EPS have been solved with the inclusiong of binaries and binary evoltuion is an important ingredient in EPS and also a subject of this symposium. ",
        "conclusions": ""
    },
    "0910/0910.5562.txt": {
        "abstract": "We monitored the BL Lac object S5 0716+714 in the optical band during October 2008, December 2008 and February 2009 with a best temporal resolution of about 5 minutes in the \\emph{BVRI} bands. Four fast flares were observed with amplitudes ranging from 0.3 to 0.75 mag. The source remained active during the whole monitoring campaign, showing microvariability in all days except for one. The overall variability amplitudes are $\\Delta$\\emph{B} $\\sim$ 0$^{m}$.89, $\\Delta$\\emph{V} $\\sim$ 0$^{m}$.80, $\\Delta$\\emph{R} $\\sim$ 0$^{m}$.73 and $\\Delta$\\emph{I} $\\sim$ 0$^{m}$.51. Typical timescales of microvariability range from 2 to 8 hours. The overall \\emph{V} - \\emph{R} color index ranges from 0.37 to 0.59. Strong bluer-when-brighter chromatism was found on internight timescales. However, different spectral behavior was found on intranight timescales. A possible time lag of $\\sim$ 11 mins between \\emph{B} and \\emph{I} bands was found on one night. The shock-in-jet model and geometric effects can be applied to explain the source's intranight behavior. ",
        "introduction": "Blazars represent an extreme subclass of active galactic nuclei. They are characterized by rapid and strong variability throughout the entire electromagnetic wavebands, high and variable polarization($>$3\\%) from radio to optical wavelengths. In the unified model of AGNs, blazars make an angle of less than $10\\,^{\\circ}$ from the line of sight (Urry \\& Padovani 1995). For low-energy peaked (or radio-selected) blazars the continuum from radio through the UV or soft X-rays is mainly contributed by synchrotron radiation, while a second hump in the spectrum at higher energies usually peaks in the GeV $\\gamma$-ray band.\u00a1\u00b1\\ Blazars exhibit variability on different timescales, from years down to hours or less (See Fan et al. 2005). Understanding blazar variability is one of the major issues of active galactic nuclei studies. There may be different behavior on different timescales. Periodicity may be found on the long term, as is the case of OJ 287 (Sillanpaa et al. 1988) and other blazars (Fan et al. 2002, 2007). The periodicity of OJ 287 can be explained by a binary black hole model, with the secondary black hole passing through the accretion disk of the primary black hole twice per revolution (Sillanpaa et. al 1988; Lehto \\& Valtonen 1996). On shorter timescales, 3C 66A was claimed to have an optical period of $\\sim$ 65 days during its bright state (Lainela et al. 1999), while Mkn 501 may have displayed a period of 23 days in high energy data (Osone 2006). Recently, analyses of X-ray data have yielded excellent evidence for a quasi-period of about an hour for 3C 273 (Espaillat et al. 2008) and very good evidence for near periods of $\\sim$ 16 days for AO 0235+164 and $\\sim$ 420 days for 1ES 2321+419 (Rani et al. 2009). This may be due to a shock wave moving along a helical path in the relativistic jet (e.g., Marscher 1996) or the unstability in the accretion disk (Fan et al. 2008) for Mkn 501. Different viewing angles may also result in different level of brightness, with the source being dimmer at larger viewing angles. BL Lacertae shows variability on both intranight and internight timescales with different spectral behaviors. It shows an intranight strongly chromatic trend and an internight midly chromatic trend, indicating two different components operating in the engine (Villata et al. 2004; Hu et al. 2006). Understanding the shortest variability timescale is of special importance. Brightness changes of up to a few tenths of a magnitude over the course of a night or less is known as intra-night optical variability (INOV) (Wagner \\& Witzel 1995) or the so-called microvariability. It can bring new insight into the understanding of the nature of blazars, probing into the innermost structure down to microparsecs, putting constraints on the size of the source and its physical environments.   Microvariability was first discovered in the sixties by Matthews and Sandage (1963), who found 3C 48 changed $0^{m}$.04 in the \\emph{V} band in 15 minutes. But their results were not taken seriously and were considered due to instrumental errors. However, with the development of CCD, microvariability was confirmed to be the intrinsic nature of active galactic nuclei, especially for blazars (e.g., Miller et al. 1989). From a complete sample of BL Lac objects (Stickel et al. 1991), Heidt and Wagner (1996) found 80\\% exhibited microvariability, proving it to be the nature of BL Lac objects. Romero et al. (1999) detected microvariability in 60\\% of their selected sample of 23 southern AGNs. Microvariability became extensively observed since the eighties. Up to now, the reasons for it are still unclear. Many different models have been proposed to explain this phenomenon. It conceivably may be due to eclipses if there is a binary black hole system (Xie et al. 2002). The interaction of relativistic shock waves and inhomogeneous jet can also cause microvariability in the jet (Maraschi et al. 1989; Qian et al. 1991; Marscher 1992; Marscher 1996). Other models like instabilities and perturbations in the accretion disk can explain some of the microvariability phenomena (e.g., Mangalam \\& Wiita 1993; Wiita 1996; Fan et al. 2008). In order to understand the radiation mechanism and the structure of the radiating region, long-term and multiwavelength observations are needed. The BL Lac object S5 0716+714 is one the brightest BL Lac objects noted for its microvariability. Its high duty cycle means that it is always in an active state (Wagner \\& Witzel 1995). It has been the target of many monitoring programs (e.g., Wagner et al. 1996; Qian et al. 2001; Raiteri et al. 2003). Montagni et al. (2006) reported the fastest variability rate of 0.1-0.12mag/hr. Five major outbursts have been observed so far, occurring at the beginning of 1995, in late 1997, in the fall of 2001, in March 2004 and at the beginning of 2007 (Raiteri et al. 2003; Foschini et al. 2006; Gupta et al. 2008). These five outbursts indicate a long-term variability timescale of $\\sim$ 3.0$\\pm$0.3 years (e.g., Raiteri et al. 2003). Correlated radio/optical variability has been observed for the source. In a 4-week monitoring program, Quirrenbach et al. (1991) discovered a period change of  1 day to 7 days in both radio and optical bands. Heidt \\& Wagner (1996) reported a period of 4 days in the optical band while Qian, Tao \\& Fan (2002) derived a 10-day period from their 5.3 yr optical monitoring. Recently Gupta et al. (2009) have found good evidence for quasi-periods in five of the 20 best quality nightly data sets of Montagni et al. (2006); these ranged from $\\sim$ 23 to $\\sim$ 75 minutes.     The spectral change of S5 0716+714 has been observed extensively (e.g., Ghisellini et al. 1997; Raiteri et al. 2003; Villata et al. 2004; Gu et al. 2006; Wu et al. 2005). Different behaviors have been reported. Raiteri et al. (2003) reported different chromatism in different timescales in their 8-year monitoring program. Sometimes the source was bluer when brighter, sometimes the opposite, sometimes no spectral change was seen despite changes in brightness. Ghisellini et al. (1997) and Wu et al. (2005) found the source exhibited a bluer-when-brighter chromatism when it was in fast flares. However, Stalin et al. (2006) found even if the source was in fast flare, there was no color change with brightness. In this paper we concentrate on the microvariability and spectral changes of S5 0716+714. We monitored the source from October 25 to 30 2008, December 23 to 29 2008 and February 3 to 10 2009. The temporal resolution was around 5 to 8 minutes in the four optical bands (\\emph{BVRI}). Because of the high temporal resolution, we can provide high quality data with accurate results.  The paper is organized as follows: Section 2 describes observations and data reduction procedures. Section 3 presents the results. Discussions are reported in Section 4. A summary is given is Section 5. ",
        "conclusions": "The microvariability of S5 0716 + 714 has been observed by many authors. Some ultra-rapid fluctuations on timescale $\\leq$ 1.0 hour were reported. Qian et al. (2002) recorded a brightness increase of $\\Delta$\\emph{V} $\\sim$ 0.78 mag in 9 mins in their 5.3-yr monitoring programme. Xie et al. (2004) reported $\\Delta$\\emph{B} $\\sim$ 0.55 mag on a timescale of 36 mins. On longer timescales, Gu et al. (2006) found the magnitude change of $\\Delta$\\emph{V} = 0.28 mag in $\\sim$ 5 hours. In our monitoring programme, no ultra-rapid fluctuations were detected. The shortest timescale was $\\sim$ 2 hours, corresponding to $\\Delta$R $\\sim$ 0.046 mag on JD 2454826 while the longest timescale were $\\sim$ 8 hours, corresponding to $\\Delta$R $\\sim$ 0.245 mag, which happened on JD 2454871 - JD 2454872. In our whole monitoring campaign, the timescales are generally a few hours with amplitude of $\\sim$ 0.04 - 0.28 mag.  Time lags between different optical bands have also been reported. Stalin et al. (2006) found possible time lags of $\\sim$ 6 and $\\sim$ 13 mins between the \\emph{V} and \\emph{R }bands in their 2 nights of observations of the source. Qian et al. (2000) reported similar results in which they found a time lag of $\\sim$ 6 mins between the \\emph{V} and \\emph{I} bands. From densely sampling data, Villata et al. (2000) presented a time lag with a strict upper limit of $\\sim$ 10 mins between the \\emph{B} and \\emph{I} band. In this work, we tried to find out the time lag between \\emph{B} band and \\emph{I} band using the discrete cross-correlation function, DCF, suggested by Edelson \\& Krolik (1988) for unequally spaced data. It was performed only on the day with high-quality data (error $\\sim$ 0.003 and $\\sim$ 0.005 for \\emph{B} and \\emph{I} band respectively) and dense sampling rate (temporal resolution $\\sim$ 5 minutes), which is JD 2454826, the day with with a hint of periodicity in its light curve. Fig. 16 presents the results of the calculations. The dashed line indicates the centroid which was calculated using the method proposed by Peterson (2001) and we found the barycenter using data points located near the peak value, DCF$_{peak}$, specifically, those greater than 0.8DCF$_{peak}$. The expression is: \\begin{equation} DCF_{centroid} ={\\frac{\\sum \\tau DCF(\\tau)}{\\sum DCF(\\tau)}} \\end{equation} We found a time lag of $\\sim$ 11 minutes, with the \\emph{B} band leading \\emph{I} band. However, given the sampling rate of $\\sim$ 5 minutes and the negligible difference between the DCF values at -20 minutes and 0 minutes, this cannot be considered a convincing measurement of a lag.  This result is consistent with Villata et al. (2000), who presented a strict upper limit of $\\sim$ 10 minutes to a possible delay between \\emph{B} - and \\emph{I} - band variations using high-quality, densely sampled data on a single night. This result is expected in the shock-in-jet model but not in most disc-based variability models (e.g. Wiita 2006).  In order to quantitatively analyze possible periods, we performed structure function (SF) on JD 2454826, the only day with a lightcurve that shows a hint of a period. SF, discussed fully by Simonetti el al. (1985), is a common tool to search for periodicities and timescales of variation . It identified a timescale of 259 mins, 287 mins, 280 mins and 263 mins and a period of 482 mins, 497 mins, 498 mins and 491 mins were found for \\emph{B}, \\emph{V}, \\emph{R} and \\emph{I} band respectively. All these results are consistent with visual inspection. The results are shown in Fig. 17. Of course, since only one ``cycle'' is present that night, the presence of a period is no more than speculative; such a possible period does not appear to extend to the previous night, for which there is also data.  Only observations made with multiple ground based telescopes at different longitudes (or space based telescopes) can convincingly find intranight periods that substantially exceed $\\sim$ 1 hr.   The spectral variability of blazars has been investigated by many authors (e.g., Ghisellini et al. 1997; Fan \\& Lin 1999; Romero et al. 2000; Raiteri et al. 2003; Villata et al. 2000, 2004). Raiteri et al. (2003) found different spectral behavior for the S5 0716 + 714 in short timescales. Sometimes the source was bluer when brighter, sometimes the opposite and sometimes no spectral change was seen. Stalin et al. (2006) found no clear evidence of color variation with brightness in either their internight or intranight monitoring of the source for a fortnight. Ghisellini et al. (1997) and Wu et al. (2005) reported a BWB trend during fast flares but this trend is insensitive in the long-term. However, Wu et al. (2007) noticed that the source is bluer when brighter on both intranight and internight timescales but this trend was not present in the long-term data.  In our monitoring campaign, the source also showed different spectral change behavior. On the long term, the source showed a bluer-when-brighter behavior. Unlike previous studies which reported consistent trends on spectral behavior on intranight timescales (e.g. Ghisellini et al. 1997, Stalin et al. 2006, Wu et al. 2005), no consistent spectral behavior is found among all those nights displaying microvariability in our present study. The source displayed BWB chromatism when it was either bright or dim, or when showing large or small variability amplitudes. Achromatism was also found when the source was in the same conditions, suggesting different mechanisms for microvariability. The BWB behavior is consistent with shock-in-jet model (Wagner et al. 1995). In this model, shocks form at the base of the jet and propagate downstream, accelerating electrons and compressing magnetic fields and resulting in the observed variability. The model predicts a BWB phenomenon and an irregular lightcurve, as observed in our cases showing this phenomenon. In the case of JD 2454825 and JD 2454826, a symmetric lightcurve was observed for the former case while a periodic lightcurve for the latter. Such regular lightcurves are rarely found. Wu et al. (2005) reported a single ``period'' of a sine-like light curve on intranight timescales for two nights during their observations of S5 0716+714. Using more sophisticated techniques Gutpa et al. (2009) found multiple oscillations to be present during 5 nights out of the 20 highest quality light curves of a total sample of 102 nights of data by Montagni et al. (2006). Unlike ours, their lightcurves are sinelike with smooth turns while ours are sawtooth-like with sharp turnoffs. The symmetric lightcurves on JD 2454825 showed a correlation between color and magnitude, with the correlation coefficient r = 0.616. So the variation may still be due to intrinsic reasons. However, for the periodic lightcurve on JD 2454826, no strong correlation between color and magnitude is found, with the correlation coefficient r = 0.150. Such achromatic lightcurves may be produced by geometric effects like microlensing or a lighthouse effect produced by different amounts of Doppler boosting induced by a helical structure to the jet (e.g. Wagner \\& Witzel 1995). Microlensing predicts a strictly symmetric lightcurve, which cannot explain the concave shape of the second halves of the light curves in our case. Therefore a lighthouse effect is the most probable mechanism for explaining periodic components in blazar lightcurves as this type of variation is likely to be achromatic (e.g., Camenzind \\& Krockenberger 1992; Gopal-Krishna \\& Wiita 1992)."
    },
    "0910/0910.0225_arXiv.txt": {
        "abstract": "{Plasma processes close to supernova remnant shocks result in the amplification of magnetic fields and in the acceleration of electrons, injecting them into the diffusive acceleration mechanism.} {The acceleration of electrons and the magnetic field amplification by the collision of two plasma clouds, each consisting of electrons and ions, at a speed of 0.5c is investigated. A quasi-parallel guiding magnetic field, a cloud density ratio of 10 and a plasma temperature of 25 keV are considered.} {A relativistic and electromagnetic particle-in-cell simulation models the plasma in two spatial dimensions employing an ion-to-electron mass ratio of 400.} {A quasi-planar shock forms at the front of the dense plasma cloud. It is mediated by a circularly left-hand polarized electromagnetic wave with an electric field component along the guiding magnetic field. Its propagation direction is close to that of the guiding field and orthogonal to the collision boundary. It has a frequency too low to be determined during the simulation time and a wavelength that equals several times the ion inertial length. These properties would be indicative of a dispersive Alfv\\'en wave close to the ion cyclotron resonance frequency of the left-handed mode, known as the ion whistler, provided that the frequency is appropriate. However, it moves with the super-alfv\\'enic plasma collision speed, suggesting that it is an Alfv\\'en precursor or a nonlinear MHD wave such as a Short Large-Amplitude Magnetic Structure (SLAMS). The growth of the magnetic amplitude of this wave to values well in excess of those of the quasi-parallel guiding field and of the filamentation modes results in a quasi-perpendicular shock. We present evidence for the instability of this mode to a four wave interaction. The waves developing upstream of the dense cloud give rise to electron acceleration ahead of the collision boundary. Energy equipartition between the ions and the electrons is established at the shock and the electrons are accelerated to relativistic speeds.} {The magnetic fields in the foreshock of supernova remnant shocks can be amplified substantially and electrons can be injected into the diffusive acceleration, if strongly magnetised plasma subshells are present in the foreshock, with velocities an order of magnitude faster than the main shell.} ",
        "introduction": "Supernova remnants (SNRs) emanate energetic electromagnetic radiation, which demonstrates the acceleration of electrons to ultrarelativistic speeds \\citep{RelEl1,Radiation} and the generation or amplification of magnetic fields \\citep{MagAmp1,MagAmp2}. The likely origin of the accelerated electrons and of the strong magnetic fields is the SNR shock \\citep{Marco1,Marco2}.  The nonrelativistic expansion speed of the main SNR blast shell \\citep{Shockspeed1,Shockspeed2} and the weak magnetic field of the ambient medium \\citep{MagAmp2}, into which this shell is expanding, are obstacles to the magnetic field amplification by plasma instabilities and to the electron acceleration out of the thermal plasma distribution to moderately relativistic energies. Such an acceleration is needed for their injection \\citep{Injection1,Injection2,Injection3,Gyrosurf} into the diffusive shock acceleration process (See \\citet{Diffus}) so that they can cross the shock transition layer repeatedly. Electrostatic instabilities dominate for nonrelativistic flows in unmagnetized plasmas \\citep{Bret,BRETAPJ} and they can neither accelerate the electrons to highly relativistic speeds \\citep{Sircombe} nor amplify the magnetic fields.  The electrons could be accelerated by plasma based charged particle accelerators \\citep{PlAc}, by electron surfing acceleration \\citep{Surfing1,Surfing2,Surfing3}, double layers \\citep{DL2,DL1} or by processes that exploit a velocity shear in the plasma outflow \\citep{Shear}. It has, however, not yet been demonstrated with multi-dimensional and self-consistent simulations that these mechanisms can indeed achieve the required electron acceleration and magnetic field amplification. The non-relativistic shocks, which are found between the main SNR blast shell and the ambient medium, can also probably not transfer significant energy from the ions to the electrons and accelerate the latter to relativistic speeds \\citep{Shock4,Shock6,Shock5,Shock3,Sorasio,Shock2,Shock7}.  A viable acceleration mechanism may develop ahead of the main SNR shock, if we find subshells that outrun the main blast shell. These subshells may move faster than the typical peak speed of 0.2c of the main shell. An expansion speed as high as 0.9c may have been observed for a subshell ejected by the supernova SN1998bw \\citep{Shockspeed1}. Most supernovae are less violent and their subshells are probably slower. The density of the subshell plasma is well below that of the main blast shell and its dynamics will be influenced to a larger extent by the upstream magnetic field than the dynamics of the latter. This is true in particular, if the upstream magnetic field has been pre-amplified by the cosmic rays \\citep{Winske,MagAmp3,MagAmp4,Pohl,Riq}. A fast magnetized shock would form in the foreshock of the main SNR shock that can result in a stronger magnetic field amplification and electron acceleration.  We examine with a particle-in-cell (PIC) simulation the formation of a shock in a plasma, in which a strong guiding magnetic field is quasi-parallel to the plasma expansion direction. Whistler waves, which become Alfv\\'en waves at low frequencies, occuring at such shocks can be efficient electron accelerators \\citep{Oblique3,Injection3,Gyrosurf,ANOTHERSLAM,Ken}. Whistlers and Alfv\\'en waves are circulary polarized if they propagate parallel to the guiding magnetic field. We briefly summarize their properties and focus on the low-frequency modes with a left-hand polarization. These modes are qualitatively similar to the quasi-parallel propagating ones we consider here. A more thorough description of the dispersion relation of two-fluid waves and the shift of the resonance frequencies for quasi-parallel propagation can be found elsewhere \\citep{Treumann}.  The dominant waves, which we will observe, grow in a plasma with an electron gyrofrequency that exceeds the plasma frequency. As we increase the frequency in such a strongly magnetized plasma towards the ion cyclotron frequency, the Alfv\\'en waves with a left-hand circular polarization become dispersive. These waves resonate with the ions and their frequencies remain below the ion cyclotron frequency. Alfv\\'en modes with a left-hand circular polarization just below the ion cyclotron frequency are called ion whistlers. Whistlers are predominantly electromagnetic for small propagation angles relative to the guiding magnetic field, as it has been discussed for high-frequency ones by \\citep{Tokar}, and if they have low wavenumbers. Ion whistlers or dispersive Alfv\\'en waves develop a field-aligned electric field component close to the resonance frequency, by which they can interact nonlinearly with the plasma particles and accelerate them \\citep{Ken}. Any wave growth will furthermore result in an increasing energy density of the magnetic field.  Alfv\\'en waves and other magnetohydrodynamic (MHD) waves can grow to amplitudes, at which they start to interact nonlinearly with the plasma \\citep{STASI}. Short Large Amplitude Magnetic Field Structures (SLAMS) are nonlinear MHD waves occuring at quasi-parallel shocks and may be relevant for our simulation. The SLAMS can be efficient electron accelerators \\citep{CME2}. Their magnetic amplitude reaches several times that of the background field and they can propagate with a super-Alfv\\'enic speed, because they convect with the ions \\citep{SLAMSPEED,SLAMS}. SLAMS have also been observed in simulations \\citep{Scholer}. The acceleration of electrons by sub-structures of quasi-parallel shocks, which may be SLAMS, has been observed in the solar corona \\citep{CME}.  The absence of self-consistent kinetic models of oblique shocks implies that they can currently be studied only numerically with particle-in-cell (PIC) \\citep{Code2,Code1} or with Vlasov simulations \\citep{Arber,Sircombe}. The pioneering PIC simulations of plasma slabs, which collide with a speed of 0.9c and in the presence of an oblique magnetic field \\citep{Oblique1,Oblique2}, have evidenced the formation of a shock that accelerated the electrons to ultrarelativistic speeds and amplified the magnetic field. A more recent PIC simulation study \\citep{Shock1} has shown, that the shock formation is triggered by an energetic electromagnetic structure (EES). The simulation could demonstrate that an approximate equipartition of the ion, electron and magnetic energy densities is established. However, these simulations could only resolve one spatial dimension due to computer constraints, which is not necessarily realistic for mildly relativistic collision speeds.  We perform here a case study with initial conditions, which are similar to those employed by \\citet{Shock1}. We reduce the collision speed to 0.5c and lower the temperature. The plasma cloud representing the subshell is ten times denser than the plasma cloud that represents the ambient plasma (interstellar medium), into which the shell expands. A guiding magnetic field is quasi-parallel to the plasma flow velocity vector and it results in an electron gyrofrequency that equals the electron plasma frequency of the dense cloud. These bulk plasma parameters have been selected with the intention to enforce a planar shock front \\citep{Nish,Mag,Shock1}, by which we can model the shock in one- or two-dimensions in space. The reduced ion-to-electron mass ratio we use allows us to model this collision in form of a 2(1/2)D particle-in-cell (PIC) simulation, which resolves two spatial and three momentum dimensions.  Our results are summarized as follows: We find higher-dimensional structures (density filaments), which initially form at the front of the tenuous plasma cloud and expand in time. The front of the dense cloud, which turns out to be the most relevant structure, remains planar and its filamentation is delayed but not suppressed by the guiding magnetic field and the high temperature. An EES grows ahead of the dense plasma cloud before it has become filamentary and the magnetic amplitude of the EES reaches a value several times the one of the initial guiding magnetic field. The EES is pushed by the dense cloud and its phase speed in the rest frame of the tenuous cloud is comparable to the cloud collision speed or twice the Alfv\\'en speed in the tenuous cloud. The front of the EES expands at an even higher speed. Its high speed and amplitude may imply that the EES is a SLAMS. The front of the dense cloud and, consequently, the EES are slowed down by the electromagnetic wave-particle interaction. It is thus not possible to define a rest frame moving with a constant speed, in which the EES is stationary. This would be necessary to measure the frequency of the EES accurately. However, the amplitude distribution of the EES suggests that its oscillation frequency is below the ion cyclotron frequency. Electromagnetic waves are destabilized by the EES ahead of the cloud overlap layer through what we think is a four-wave interaction \\citep{Instability}. The electrons are accelerated in the combined wave fields of these waves and in the forming shock to highly relativistic speeds. The simulation shows though that the strongest electron acceleration occurs at the position, where the shock-reflected ion beam is forming. The plasma collision results in a substantial electron acceleration and also in an amplification of the magnetic energy density by one order of magnitude within the EES and the forming shock. Both values are probably limited by the reduced ion mass of 400 electron masses, which we must employ. Radiative processes, which are not resolved by the PIC code, will at this stage start to influence the shock evolution \\citep{Schl1,Schl2} and we stop the simulation.  This paper is structured as follows. Section 2 discusses the particle-in-cell simulation method and the initial conditions. Section 3 presents our results, which are discussed in more detail in Section 4. ",
        "conclusions": "We have described in this paper the collision of two plasma clouds at the speed c/2. The ion to electron mass ratio of 400 has allowed us to model the collision in two spatial dimensions until the shock forms. Then we had to stop the simulation. The acceleration of electrons to speeds $\\sim c$ and the rapid expansion speed of the energetic electromagnetic structure (EES) imply, that both will quickly reach the boundaries; hence our periodic boundary conditions become invalid. Open boundaries would allow the electrons and the wave energy to flow out of the system. However, the instabilities driven by these beams \\citep{Martins} and by the EES are important for the upstream dynamics and the latter will be adversely affected by open boundaries.  The filamentary structures can form and merge in the 2D geometry we consider here up to the instant when the magnetic repulsion of two filaments with oppositely directed current enforces their spatial separation \\citep{Davidson,Lee}, at least in the absence of an oblique magnetic field. Then no further mergers occur since only one dimension is available orthogonal to the flow velocity vector. A realistic 3D PIC simulation would allow the filaments to move around each other and merge with other filaments of equal polarity \\citep{Lee}. A 3D PIC simulation is, however, currently impossible for our case study involving ions, because of the computational cost involved in resolving ion and electron scales. Simulations in three spatial dimensions are now feasible for the case of leptonic shocks \\citep{Nishi2}.  The densities of the clouds differ by a factor of 10 and the collision is asymmetric, as in the simulation by \\citet{Oblique1,Oblique2,Sorasio}. Initially the magnetic field is uniform and quasi-parallel to the flow velocity vector. Its significant strength together with the high plasma temperature of 25 keV and the unequal cloud densities reduce the growth rate of the filamentary instabilities \\citep{Mag,Bellido}.  \\citet{Shock1} have previously probed the higher $\\gamma$ regime appropriate for the internal shocks of gamma-ray bursts. Here we consider the mildly relativistic regime, with a collision speed $0.5c$ between both clouds. Such a speed might be realistic for a plasma subshell outrunning the main SNR shock. Such subshells can reach relativistic flow speeds for particularly violent SNR explosions \\citep{Shockspeed1}. In our initial conditions, the magnetic energy surpasses  the thermal energy of the dense plasma slab by a factor of 5. The plasma flow implies, however, that the box-averaged plasma kinetic energy density exceeds the magnetic energy density by an order of magnitude. We summarize several aspects of our results.  \\subsection{Effects due to initial conditions}  Our initial conditions have resulted in the formation of planar wave and plasma structures, the most important one being the EES. We think that the EES grew out of a localized seed magnetic field pulse driven by the spatial gradient of the convection electric field. The plasma upstream of the dense cloud is destabilized by this electromagnetic structure and the EES expands at the speed $0.87c$ in the reference frame of the tenuous cloud. The energy for its growth and expansion is provided by the kinetic energy of the upstream medium, which moves with respect to the EES. The shock speed $\\la v_c$ then implies that we have a coherent magnetic layer that expands its width at a speed of at least $(0.87c-v_c)/(1-0.87 v_c / c) \\approx 0.65c$, measured in the reference frame of the tenuous cloud. It covered about 80 ion skin depths at the end of the simulation, showing no signs of a slowdown.  The EES is a consequence of our initial conditions and the growth of the seed magnetic field amplitude could probably be reduced but not suppressed by a smoother change of the convection electric field, which can be achieved by a gradual change of the plasma convection speed \\citep{Oblique1,Oblique2}. However, the seed magnetic field could be provided also by waves with a short wavelength, e.g. whistlers, and it is thus not unphysical.  Structures with strong magnetic fields, similar to the EES and known as SLAMS, are frequently observed close to quasi-parallel shocks in the solar system plasma and they can accelerate electrons to high energies. They are thus potentially important also for SNR shock physics. Our initial conditions provide a possibility to let nonlinear MHD waves grow out of a simple simulation setup for a further study. The EES is moving with the ions of the dense cloud and it modulates the ions and electrons of the tenuous cloud, thereby gaining energy. Its growth probably requires an asymmetric plasma collision.  \\subsection{Shock formation}  In this paper we modelled the formation of the shock from the initial collision of two plasma clouds. The signatures of the shock are evident, including visible thermal broadening behind the shock and a dense shock ramp. While filamentation structures form ahead of and behind the shock, we note that the structure is basically planar in the critical foreshock area, where electron acceleration is expected to occur. This means that one-dimensional simulations will be relevant in this region, allowing much higher resolution, increased particle number (resulting in lower particle noise and a better phase space resolution) and a higher ion-electron mass ratio, than is currently found in two and three-dimensional simulations.  It is evident from the simulation that the filamentation is not fully suppressed by the guiding magnetic field and by the high plasma temperature. Its amplitude has been set such that it should suppress the electron filamentation if the plasma would be spatially uniform \\citep{Mag}. This amplitude is apparently insufficient to suppress the slower filamentation of the ion beams and we could even see evidence for an electron beam filamentation just behind the front. The front of the dense plasma cloud maintains its planarity throughout the simulation, but even here the density was non-uniform along the boundary. The onset of the filamentation was, however, delayed. The likely cause is the high density gradient across the front, which alters the electron and ion skin depths and thus the characteristic scale of the filaments. The gradient is caused by the slowdown of the ions by the magnetic field of the EES and by the electron acceleration. The electrons are confined at the front in the direction of the shock normal so that they preserve the quasi-neutrality of the plasma, but they can move orthogonally to it. The latter results in a drift current.  \\subsection{Field amplification}  \\citet{MagAmp2} and \\citet{MagAmp1} have described observations of magnetic field amplification above the value expected from shock compression in SNRs. At the final simulation time the magnetic field energy density is increased in strength by over one order of magnitude, exceeding by far that expected from the magnetic field compression by the shock. A shock compresses only the magnetic field component perpendicular to the shock normal, which is weak in our case, and the amplification of its energy density can reach a factor of 4-7. The magnetic energy density at the simulation's end has been comparable to the box-averaged kinetic energy density in an interval spanning about 10-20 ion skin depths. The magnetic energy density due to the EES was at least twice as high as that of the background field in an interval covering 50 ion skin depths. Even if we take into account that the kinetic energy density close to the shock is increased by the accumulation of plasma, the magnetic energy density still constitutes a sizeable fraction of the local total plasma kinetic energy density. The EES has provided the main contribution to the magnetic energy density and exceeded that due to the filaments downstream by two orders of magnitude.  Throughout this paper, we used normalized quantities in our simulation and we can scale the magnetic field amplitude to the relevant plasma conditions. If we set the electron density of the dense cloud to 1 cm$^{-3}$, we would obtain a peak magnetic field with a strength of 10 mG. However, we have to point out that our initial magnetic field amplitude has been higher than that expected for the ambient plasma, even if we take into account its amplification by cosmic ray-driven instabilities, and our simulation results may not be directly applicable.  Where does the extra field come from? Amplification of the magnetic field can occur from the electron drift current arising from the $\\vec{E} \\times \\vec{B}$ drift motion in a layer close to the shock that is narrower than the ion gyroradius but wider than the electron gyroradius, see, e.g. \\citet{BaumjohannTreumann}. The current adds to the shock current and increases the jump in the perpendicular magnetic field. This can only occur when the ion and electron gyroradii differ, i.e. not in a pair plasma. The EES has a significant $E_x$-component and $|\\vec{B}_\\perp| \\approx |\\vec{B}_0|$. We thus obtain a $\\vec{E}\\times \\vec{B}$ drift orthogonal to the flow velocity vector.  We have also found that the requirement to maintain quasi-neutrality of the plasma implies that the upstream electrons are dragged with the upstream ions across the EES, which moves with the shock. The resulting $\\vec{v}\\times\\vec{B}$ drift accelerates the electrons orthogonally to the shock propagation direction, further enhancing the net current and the magnetic field. Finally, magnetic fields of SLAMS are provided by the current due to the gyro-bunched ions, which rotate in the plane orthogonal to the wavevector.  These mechanisms increase the mean magnetic field, and are different from the instability described by \\citet{MagAmp3} who has described a cosmic ray streaming instability which can amplify turbulent magnetic fields ahead of the shock. We can exclude Bell's instability here since we have not found energetic particles with a significant density moving upstream, which would provide the net current that drives this instability.  \\subsection{Electron acceleration and upstream wave spectrum}  The shock retains its planar structure after it forms. A circularly polarised large-scale precursor wave, the EES, expands into the foreshock. Its wavelength is several times the ion skin depth. It gradually rotates the quasi-parallel magnetic field into a quasi-perpendicular one at the current layer and it forces the incoming ions and electrons to interact with it nonlinearly. The ions are gyro-bunched and some of the incoming ions of the tenuous cloud are reflected by the forming shock. We have found evidence of a parametric instability \\citep{Instability} of the EES ahead of the foreshock and we could find at least two waves that may be the result of this parametric decay. These waves appear at late times, when the EES has expanded in space and is thus sufficiently monochromatic. They grow to an amplitude that introduces oscillations of the mean velocity of the ions of up to $c/5$.  The interplay of the short-scale charge density waves and magnetowaves causes the electrons to be accelerated to highly relativistic speeds upstream of the forming shock. Similar electron acceleration (injection) mechanisms upstream of shocks involving whistler waves have been proposed by \\citet{Oblique3,Injection3,Gyrosurf}. The strongest electron acceleration is, however, observed at the location where the shock-reflected ion beam is developing. The electrons are accelerated to a peak Lorentz factor of 120 and their energy gain is thus comparable to the energy associated with the velocity change of the shock-reflected ions. If the electron acceleration is accomplished by the electromagnetic fields that reflect the incoming upstream ions, then the energy gain of the electrons may scale with the ion mass. We may expect in this case that the electrons are accelerated to a Lorentz factor $\\gamma_M \\approx 120 m_p / m_i$ that is $\\gamma_M \\approx 550$ if we would use the correct proton to electron mass ratio.  It is interesting to see if this type of electron acceleration can also occur close to Solar system shocks. Let us consider the Earth bow shock as one of the best known collisionless shocks and let us assume a Solar wind speed of $4 \\times 10^5$ m/s to $7.5 \\times 10^5$ m/s. A specular reflection of the incoming Solar wind protons by the bow shock would change their energy by about 0.8 keV to 3 keV. Electrons with such energies are observed in a thin sheet close to the shock surface of perpendicular shocks \\citep{BowShock}. Perpendicular shocks are capable to produce shock-reflected ion beams \\citep{Shock6} and the electron acceleration mechanism we observe here may work also at the Earth bow shock.  \\subsection{Future work}  This simulation study was concerned with the collision of two plasma clouds at a mildly relativistic speed. Its purpose has been to better understand the conditions and the mechanisms involved in the formation of a shock. This shock will move at an essentially nonrelativistic speed below the initial collision speed and, thus, be relevant for fast SNR flows. A two-dimensional simulation geometry was necessary to assess the importance of the multi-dimensional filamentation instability for the shock dynamics. A quasi-parallel guiding magnetic field was used to slow down this filamentation, resulting in a planar (one-dimensional) shock.  The formation of the shock could be observed, but the simulation had to be stopped at this time due to computational constraints. This simulation has, however, revealed several aspects that should be examined in more detail in more specialised simulation studies.  The EES probably formed in response to our initial conditions. i.e. the sharp jump in the convection electric field at the cloud collision boundary. It has to be investigated if the EES also forms if this jump is decreased, for example by a higher-order field interpolation scheme or by a gradual decrease of the convection electric field by a smooth change in the plasma convection speed. The simulations by \\citet{Oblique1,Oblique2} suggest that this will leave unchanged the magnetic field amplification and the electron acceleration. However, the EES may not be so strong and coherent.  The magnetic field amplitude in the present study is higher than it is realistic for a SNR scenario. Future studies must address how the shock formation depends on lower amplitudes of the guiding magnetic field. Computationally inexpensive parametric simulation studies that resolve only one spatial dimension may provide insight. Initial studies not discussed here indicate that the shock formation is delayed by a decreasing magnetic field amplitude.  It is also necessary to follow the plasma collision for a longer time. An important aspect is here how far the EES can expand upstream and how strong the downstream magnetic field is. The planarity of the EES and of the shock boundary may permit us to use one-dimensional simulations, by which we can expand by at least an order of magnitude the box size along the collision direction.  An one-dimensional simulation also allows us to examine with a larger number of particles per cell and, thus, lower noise levels the secondary instabilities driven by the EES. We have found evidence for an instability of the EES to a four-wave interaction. Lower noise levels would allow us to compare the amplitudes and phases of the EES with those of the secondary waves, which is necessary to demonstrate a coherent wave interaction. Extending the simulation time together with suitable initial conditions would, potentially, allow us to investigate what happens if the speed of the EES decreases below the local Alfv\\'en speed. It is possible that the EES decouples from the shock and propagates independently in form of an Alfv\\'en wave packet.  Finally, it would be interesting to reduce the collision speed to about $c/10$, which is close to the expansion speed of the SNR shock, to see if and how many electrons are accelerated to relativistic speeds. This will provide insight into the electron injection efficiency of oblique shocks and, thus, into the ability of SNR shocks to accelerate electrons to cosmic ray energies."
    },
    "0910/0910.5400_arXiv.txt": {
        "abstract": "{} {Recent studies have detected linear polarization in L dwarfs in the optical I band. Theoretical models have  been developed to explain this polarization. These models predict higher polarization at shorter wavelengths. We discuss the polarization in the R and I band of 4 ultra cool dwarfs.} {We report linear polarization measurements of 4 ultra cool dwarfs in the R and I bands using the Intermediate dispersion Spectrograph and Imaging System (ISIS) mounted on the 4.2m William Herschel Telescope (WHT). } {As predicted by theoretical models,  we find a higher degree of polarization in the R band when compared to polarization in the I band for  3/4 of these ultra cool dwarfs. This suggests that dust scattering asymmetry is caused by oblateness . We also show how these measurements fit the theoretical  models. A case for variability of linear polarization is found, which suggests the presence of randomly distributed dust clouds. We also discuss one case for the presence of a cold debris disk. } {} ",
        "introduction": "A large number of ultra cool dwarfs have been detected in the last decade, and our understanding of these faint objects has kept improving. One of the challenging and fundamental aspects in the study of these objects is to understand the properties and distribution of condensate dust in the atmosphere. Observations of L dwarfs with effective temperatures of 1400-2200 K have led to the investigation of dust condensates in their atmospheres (\\cite {Kirkpatrick}; \\cite {Tsuji1996}). Because of complete gravitational settling, grains are expected to condense beyond the visible atmosphere for objects with effective temperatures below 1400 K (T-Dwarfs -  \\cite{Allard 2001}; \\cite{Chabrier00}). At higher effective temperatures (1400-2200 K), grains can be present in the visible atmosphere because of incomplete gravitational settling (\\cite {Burrows and Sharp}; \\cite {Burrows 2001}; \\cite {Ackerman 2001}; \\cite {Allard 2001};  \\cite{Tsuji 2004};   \\cite{co03}; \\cite{helling}  ). Recent discoveries of blue L dwarfs and L-T transition type dwarfs (as identified in \\cite {Knapp  2004}; \\cite {Chiu 2006}; \\cite{Tsuji03}) have brought forth models which could explain this phenomenon (eg. \\cite {BSH 2006}, \\cite {Knapp  2004}) by mechanisms which involve dust settling. It would be very important to validate these mechanisms.  Linear polarization could be a very useful tool in understanding the  observationally poorly constrained dust properties in the atmospheres of L dwarfs. The possibility of detecting polarization at optical wavelengths from grains in the atmospheres of L dwarfs was first raised by \\cite{SKR 2001}. Fast rotation of L dwarfs  will induce the shape of their photosphere into the form of an oblate ellipsoid,(\\cite {Basri 2000}) and this nonsphericity will lead to the incomplete cancellation of the polarization from different areas of the stellar surface (\\cite {SKR 2001}). This prediction was first confirmed by the detection of linear polarization at 768 nm from a few L dwarfs by \\cite {Menard 2002}. Recently, \\cite{Zap 2005} have reported R and I band detection of linear polarization from several L dwarfs. Since polarization in the optical is unlikely to be due to Zeeman splitting of atomic or molecular lines or by synchrotron radiation, the observed  polarization can be explained by single dust scattering in a rotationally induced oblate atmosphere (\\cite {S 2003}; \\cite{SK 2005}) or it could be due to large and randomly distributed dust clouds (\\cite{Menard 2002}).  In this paper, we report polarization measurements of  3 L dwarfs (L0-L5) and one M9.5 dwarf  with WHT/ISIS in both I and R bands .  We also discuss our results comparing them with the recently published results of \\cite{goldman09}.  Our measurements show the general trend that  polarization is higher in the R band than  the one in the I band. This trend strongly supports the presence of dust in the atmosphere of L dwarfs as it is very  unlikely that any other mechanisms (such as the presence of magnetic field) can explain this observation   at optical wavelengths (\\cite{Menard 2002}). We also discuss how the theoretical models  (see Sect. 3)  successfully fit our measured data. ",
        "conclusions": "\\begin{enumerate} \\item We report linear polarizaion measurements of 4 very nearby ultra cool dwarfs in the R and I bands. \\item We find that there is a trend (3 out of 4) of a higher degree of  polarization at shorter wavelengths (R band) when compared to the I band as predicted by the theoretical models of \\cite{SK 2005}. \\item  The L0 dwarf 2MASS J17312974+2721233 is interesting because of its relatively high polarization and requires follow-up studies. \\item We also fit theoretical models to predict the dust grain size and rotational velocities of three of the ultra cool dwarfs. \\item We find evidence for variability in the linear polarization for (2MASSW J1507476-162738). This suggests atmospheric activities like dynamical variations of the cloud cover in this object. \\end{enumerate}"
    },
    "0910/0910.5546_arXiv.txt": {
        "abstract": "We report the detection of pulsations at 552~Hz in the rising phase of two type-I (thermonuclear) X-ray bursts observed from the accreting neutron star EXO~0748$-$676 in 2007 January and December, by the {\\it Rossi X-ray Timing Explorer}. % The fractional amplitude % was % 15\\% (rms). % The dynamic power density spectrum for each burst revealed an increase in frequency of $\\approx1$--2~Hz while the oscillation was present. The frequency drift, the high significance of the detections and the almost identical signal frequencies measured in two bursts separated by 11 months, confirms this signal as a burst oscillation similar to those found in 13 other sources to date. We thus conclude that the spin frequency in \\src\\ is within a few Hz of 552~Hz, rather than 45~Hz as was suggested from an earlier signal detection by \\cite{villarreal04}. Consequently, Doppler broadening must significantly affect spectral features arising from the neutron star surface, so that the narrow absorption features previously reported from an {\\it XMM-Newton}\\/ spectrum could not have arisen there. The origin of both the previously reported 45~Hz oscillation and the X-ray absorption lines is now uncertain. ",
        "introduction": "Neutron stars in low-mass X-ray binaries (LMXBs) provide observational evidence for their rapid spins reluctantly, and via increasingly diverse phenomena. Highly coherent burst oscillations, occuring only around the peak of thermonuclear (type-I) bursts, were the first such phenomenon to be discovered, and since have been detected in $\\approx14$ sources \\cite[e.g.][]{watts08}. Continuous pulsations in the persistent emission occur even more infrequently, and occur at just above the burst oscillation frequency in those sources which exhibit both \\cite[]{chak03a}. This result supports the hypothesis that the burst oscillation frequency traces the neutron star spin. Most recently, intermittent persistent pulsations have been detected in several sources, including one previously known burst oscillation source \\cite[e.g.][]{gal07a,altamirano07,casella07}.  It is uncertain why some sources show pulsations or burst oscillations and others do not. Additionally, sources which exhibit burst oscillations do not do so in every burst. In fact, the presence of oscillations can be as rare as 1 burst in 14 \\cite[for 4U~1916$-$053; see][]{bcatalog}. The mechanism by which the burst oscillations are produced is yet another uncertainty. Although oscillations are observed at high (fractional) amplitudes early in some bursts, consistent with a spreading ``hot spot'' model \\cite[e.g.][]{stroh97b}, the oscillations in the burst tail, when the burning must have spread to the entire stellar surface, are harder to explain. The oscillations may instead (or also) arise from anisotropies in the surface brightness originating from hydrodynamic instabilities \\cite[]{slu02} or modes excited in the neutron star ocean (e.g. \\citealt{cb00}; see also \\citealt{heyl04,pb05}).  The low-mass X-ray binary \\src\\ is particularly well-studied. This transient was discovered during {\\it EXOSAT}\\/ observations in 1985 \\cite[]{parmar86}, which also revealed the first thermonuclear bursts from the source as well as X-ray dipping activity. Synchronous X-ray and optical eclipses \\cite[]{crampton86} are observed once every 3.82~hr orbit. A variety of low- and high-frequency variability has been characterised \\cite[e.g.][]{homan99,homan00}, notably including a 695~Hz quasi-periodic oscillation. More recently, absorption features in the summed X-ray spectra of bursts observed by {\\it XMM-Newton}\\/ were identified as redshifted lines from near the neutron star surface \\cite[$z=0.35$;][]{cott02}. The narrowness of these features requires that the neutron star is rotating slowly, as rotation speeds $\\ga100$~Hz will broaden the line profiles to the point where they are undetectable \\cite[]{ozel03,chang06,bml06}. Subsequently, \\cite{villarreal04} detected a 45~Hz peak in the summed power spectrum of 38 thermonuclear bursts, which they interpreted as a spin frequency sufficiently slow to give negligible broadening. However, subsequent followup studies have failed to confirm the spectral line detection \\cite[e.g.][]{cott08}.  Here we present analysis of recently observed bursts from \\src\\ which suggest that the neutron star spin is not 45~Hz, but 552~Hz. ",
        "conclusions": "\\label{sec2}  There are substantial differences in the characteristics of the two burst oscillations (45~Hz and 552~Hz) now detected in \\src. The 45~Hz signal had a fractional amplitude of $\\approx3$\\% (rms) and was detected in the summed power spectrum of 38 bursts, rebinned by a factor of (typically) 64 or 128 to give a resolution of 1~Hz, calculated from light curves selecting photons within the energy range 6--60~keV, and covering intervals of typically 64 or 128~s of the burst decay. In contrast, the 552~Hz signal has a fractional amplitude of 15\\% (rms), was detected separately in unbinned power spectra calculated from 1- and 2-s light curves extracted over the full PCA energy range ($\\approx2$--60~keV) during the rise of two individual bursts, separated by 11~months. Additionally, the dynamic power density spectra give evidence for frequency evolution while the 552~Hz signal was present, similar to that observed in other burst oscillation sources. The key question is, which of these two signals traces the neutron star spin?  We consider three alternatives. First, assuming both signals are genuine, the 45~Hz signal may arise from the neutron star spin, in which case the 552~Hz signal may arise from a high-order ($m\\approx12$--13) radial mode (\\citealt{cb00}; see also \\citealt{heyl04,pb05}). However, the appearance of this mode alone is difficult to explain; one would expect the excitation of many different modes, rather than just two with widely-separated orders. Second, if both signals are real but the 552~Hz signal instead indicates the neutron star spin, there is no known mechanism that can give rise to a 45~Hz signal from the neutron star surface. As suggested by \\cite{balman09}, the 45~Hz oscillation may instead arise in the boundary layer between the disk and neutron star.  Third, the lack of a detection of the 45~Hz signal in the larger sample of bursts detected since 2004 suggests that the 45~Hz signal may have arisen from statistical fluctuations. Our analysis shows that including many more bursts in the power spectral sum has the effect only of reducing the detection significance to well below any conservative detection threshold. The signal power is also quite sensitive to other parameters of the data selection. Thus, while it is possible that the 45~Hz oscillation is transient and has not been detectable since, it is also difficult to rule out an origin in statistical noise with any confidence.  The properties of the 552~Hz oscillation closely resemble those of the burst oscillations observed in other systems, that are believed to trace the neutron star spin to within a few Hz.  In contrast, the properties of the 45~Hz oscillation are quite distinct, especially its low frequency and the inability to detect it in individual bursts. Thus, we conclude that the 552~Hz signal almost certainly traces the neutron star spin in \\src; since the signal appears to increase by up to 2~Hz during the burst rise, we expect that the true spin frequency may be a few Hz higher.  The energy dependence of the 552~Hz oscillation amplitude in \\src\\ was similar to that of other burst oscillation sources \\cite[]{muno03c}, although the overall amplitudes were higher. The suggestion of the hard ($\\ga8$~keV) photons arriving earlier than the soft photons is however contrary to previous observations. In that respect we note that the previous analysis focussed principally on oscillations present throughout the tails of bursts, whereas the oscillations in \\src\\ were present only in the rise. It is not known whether the two types of oscillations have characteristically different variation of pulse arrival time with energy.  The duty cycle of the 552~Hz oscillation (the number of bursts in which it was detected compared to the total number observed) is extremely small at $\\approx1.2$\\%, the smallest yet for all the burst oscillation sources (e.g. G08). The unprecedented scarcity of the oscillation raises the question of what was so unusual about the bursts that exhibited them. Compared to the entire sample of bursts observed by \\xte\\/ from \\src, the bursts on January 14 and December 13 had somewhat shorter timescales (calculated as the ratio of peak flux to fluence) $\\tau=12$--13~s, while the typical range is 15--30~s. Similarly, the rises were of shorter duration than average, at 2--4~s. Shorter burst timescales suggest a smaller proportion of hydrogen in the burst fuel than usual. Correspondingly, while the fluences were rather typical, the bursts reached higher than average peak fluxes. However, other bursts were observed with similar properties which did not exibit oscillations, so these properties do not uniquely determine their observeability. The January 14 burst occurred only 11.4~min after the previous event, and had a fluence only about 77\\% lower. Such short recurrence time bursts are common for \\src, and many examples have been detected by \\xte\\/ (e.g. G08) as well as {\\it XMM-Newton}\\/ \\cite[]{boirin07a}. However, the second burst in these pairs is typically much fainter than the first. Interestingly, it is also possible that the December 13 burst was the second in a closely-spaced pair, as a data gap (due to the 90~min satellite orbit) prevented observations until approximately 7~min before the event. Again, timing analysis of other such examples did not reveal any additional oscillations, however.  We also compared the properties of the persistent emission at the time when the bursts with oscillations occurred, to other public \\xte\\/ observations of the source (from G08). The persistent flux in both observations was $\\approx3\\times10^{-10}\\ \\epcs$ (2.5--25~keV), approximately equal to the 50th percentile value of the flux distribution. Similarly, the hard and soft spectral colors\\footnote{Defined as in G08 as the ratio of the background-subtracted detector counts in the (8.6--18.0)/(5.0--8.6)~keV and the (3.6--5.0)/(2.2--3.6)~keV energy bands, respectively} were in the middle of the observed range, at $\\approx0.8$ and $\\approx1.6$, respectively. Thus, we found no evidence for an unusual spectral or intensity state at the time of the bursts with oscillations.  A 552~Hz spin frequency for \\src\\ means that spectral features arising from the neutron star surface will be significantly Doppler broadened, as well as reducing the central line depth to only $\\la5$\\% of the continuum level \\cite[e.g.][]{cbw05}. Narrow lines from such rapidly-rotating objects can only arise if the system is viewed at extremely low inclinations; however, the eclipses and dipping activity in \\src\\ unambiguously indicate a high system inclination. Thus, the narrow ($\\approx0.1$~\\AA) spectral features reported by \\cite{cott02} could not arise from the neutron star surface. This conclusion is corroborated by other recent work, including the absence of lines in a deeper {\\it XMM-Newton}\\/ spectrum \\cite[]{cott08}, as well as difficulties explaining the inferred Fe column \\cite[]{cbw05}. However, the narrow spectral features are difficult to identify otherwise \\cite[e.g.][]{kong07}, and a satisfactory explanation remains elusive."
    },
    "0910/0910.4540_arXiv.txt": {
        "abstract": "Following the recent discovery of $\\gamma$ rays from the radio-loud narrow-line Seyfert 1 galaxy PMN J0948+0022 ($z=0.5846$), we started a multiwavelength campaign from radio to $\\gamma$ rays, which was carried out between the end of March and the beginning of July 2009. The source displayed activity at all the observed wavelengths: a general decreasing trend from optical to $\\gamma-$ray frequencies was followed by an increase of radio emission after less than two months from the peak of the $\\gamma-$ray emission. The largest flux change, about a factor of about 4, occurred in the X-ray band. The smallest was at ultraviolet and near-infrared frequencies, where the rate of the detected photons dropped by a factor $1.6-1.9$. At optical wavelengths, where the sampling rate was the highest, it was possible to observe day-scale variability, with flux variations up to a factor of about 3. The behavior of PMN J0948+0022 observed in this campaign and the calculated power carried out by its jet in the form of protons, electrons, radiation and magnetic field are quite similar to that of blazars, specifically of flat-spectrum radio quasars. These results confirm the idea that radio-loud narrow-line Seyfert 1 galaxies host relativistic jets with power similar to that of average blazars. ",
        "introduction": "The recent detection by \\emph{Fermi Gamma-ray Space Telescope} of $\\gamma$ rays from the radio-loud narrow-line Seyfert 1 galaxy (RL-NLS1) PMN J0948+0022\\footnote{We note that the absolute magnitude of this source is $M_B=-23.6$, so formally matches also the definition of quasars.} ($z=0.5846$) opened new and interesting questions on the unified model of active galactic nuclei (AGN), the development of relativistic jets and the evolution of radio-loud AGN (Abdo et al. 2009a, Foschini et al. 2009a). Indeed, before \\emph{Fermi}/LAT (Large Area Telescope) it was known that $\\gamma$ rays from AGN are produced in blazars and radio galaxies, but we have to add also RL-NLS1s.  NLS1s are active nuclei similar to Seyferts, where the optical permitted lines emitted from the broad-line region (BLR) are narrower than usual, with FWHM(H$\\beta$)$<2000$~km~s$^{-1}$ (see Pogge 2000, for a review). Other characteristics are [OIII]/H$\\beta<3$ and a bump of FeII, making them a peculiar class of AGN. NLS1s are different from Seyfert 2s, whose optical spectra typically display FWHM(H$\\beta$)$<1000$~km~s$^{-1}$, [OIII]/H$\\beta>3$ and no bump of FeII. NLS1s are also different from the naked AGN discovered by Hawkins (2004), a peculiar class of Seyferts without the BLR, which have [OIII]/H$\\beta>>3$. Indeed, NLS1s do have both BLR and the narrow-line region (NLR), but the BLR emits only permitted lines narrower than in Seyfert 1s (Rodr\\'iguez-Ardila et al. 2000).  NLS1s are generally radio-quiet, but a small fraction of them ($<7$\\%, according to Komossa et al. 2006), are radio-loud. It is not clear how these sources fit into the framework of radio-loud AGN. Some studies of the average non-simultaneous multiwavelength properties (from radio to X-rays) of RL-NLS1s suggested some possibilities. Komossa et al. (2006) argued that RL-NLS1s could be some young stage of quasars, while Yuan et al. (2008) found some similarities to TeV BL Lacs, but having strong emission lines they would represent the so-called ``high-frequency peaked flat-spectrum radio quasars'' conjectured by Padovani (2007). Foschini et al. (2009b) found instead that there is no one-to-one correlation of RL-NLS1s properties with any specific type of blazar or radio galaxy. In some cases, there are similarities with flat-spectrum radio quasars, while others are like BL Lacs.  Now, the first detection by \\emph{Fermi}/LAT of $\\gamma$ rays from one RL-NLS1 - namely PMN J0948+0022 - sets the definitive confirmation of the presence of a relativistic jet in these sources. The discovery of $\\gamma-$ray emission from other sources of this type (Abdo et al., in preparation) raise RL-NLS1s to the rank of $\\gamma-$ray emitting AGN. However, any average spectral energy distribution (SED) of a $\\gamma-$ray loud AGN leaves open several important questions on the mechanisms of radiation emission, such as whether the synchrotron self-Compton (SSC) or the external Compton (EC) production mechanism is dominant at high-energies and where the zone is where most of the dissipation occurs. Because PMN J0948+0022 is the first object of this new class of $\\gamma-$ray AGN, it is important to observe it for a long time, in order to understand if there is something unexpected and if its behavior is very different from blazars and radio galaxies or not.  With these aims in mind, we decided to set up a multiwavelength campaign on this source. The campaign involved several space and ground-based facilities across the whole electromagnetic spectrum, from radio to $\\gamma$ rays (in alphabetical order): ATOM (Landessternwarte), F-GAMMA (Effelsberg), e-VLBI (EVN, LBA), \\emph{Fermi}, G. Haro Telescope (INAOE), Mets\\\"ahovi, OVRO, RATAN-600, \\emph{Swift}, SMARTS, MOJAVE (VLBA), WIRO. The period covered was between 2009 March 24 and July 5. We measured variability at multiple wavebands, modelled the resulting SEDs, and compared the results to those for more typical $\\gamma-$ray blazars in the FSRQ and BL Lac classes.  Throughout this work, we adopted a $\\Lambda$CDM cosmology from the most recent \\emph{WMAP} results, which give the following values for the cosmological parameters: $h = 0.71$, $\\Omega_m = 0.27$, $\\Omega_\\Lambda = 0.73$ and with the Hubble-Lema\\^{i}tre constant $H_0=100h$ km s$^{-1}$ Mpc$^{-1}$ (Komatsu et al. 2009). ",
        "conclusions": "We thus confirm that PMN J0948+0022 -- despite being a radio-loud narrow-line Seyfert 1 -- hosts a relativistic $\\gamma-$ray emitting jet, similar to those of FSRQs, and confirms all the hypotheses adopted to model the non-simultaneous SED in Abdo et al. (2009a). This type of source can develop a relativistic jet like blazars and radio galaxies, even though the conditions of the environment close to its central spacetime singularity are quite different. This is indeed a new class of $\\gamma-$ray emitting AGN.  We have shown that the variability at multiple wavebands and the physical parameters resulting from modelling the SEDs are typical of a source midway between FSRQs and BL Lacs. The calculated powers carried by the various components of the jet are low compared to the distributions of values for FSRQ, but above those of BL Lacs (cf Celotti \\& Ghisellini 2008, Ghisellini et al. 2009), and therefore within the average range of blazar powers, despite the relatively low mass of its black hole, $1.5\\times 10^{8}M_{\\odot}$ (Abdo et al. 2009a). The $\\gamma-$ray observations performed to date have not revealed very high fluxes, i.e. above the usual threshold adopted to define an outburst in normal blazars ($F_{E>100\\rm MeV}>10^{-6}$~ph~cm$^{-2}$~s$^{-1}$). However, it is not clear yet if this is due to the duty cycle of this source -- and hence if we have just observed a minor event -- or if the different environmental conditions in the core of RL-NLS1s hampers the development of a high power jet. This question will likely be answered by the continuous monitoring that \\emph{Fermi}/LAT is performing on this and other sources of this type."
    },
    "0910/0910.4892_arXiv.txt": {
        "abstract": "{} { We used VLT/VIMOS images in the $V$ band to obtain light curves of extrasolar planetary transits OGLE-TR-111 and OGLE-TR-113, and candidate planetary transits: OGLE-TR-82, OGLE-TR-86, OGLE-TR-91, OGLE-TR-106, OGLE-TR-109, OGLE-TR-110, OGLE-TR-159, OGLE-TR-167, OGLE-TR-170, OGLE-TR-171. } { Using difference imaging photometry, we were able to achieve millimagnitude errors in the individual data points. We present the analysis of the data and the light curves, by measuring transit amplitudes and ephemerides, and by calculating geometrical parameters for some of the systems\\thanks{Photometry of the transiting objects is available at the CDS via anonymous ftp to cdsarc.u-strasbg.fr}. } { We observed 9~OGLE objects at the predicted transit moments. Two other transits were shifted in time by a few hours. For another seven objects we expected to observe transits during the VIMOS run, but they were not detected. } { The stars OGLE-TR-111 and OGLE-TR-113 are probably the only OGLE objects in the observed sample to host planets, with the other objects being very likely eclipsing binaries or multiple systems. In this paper we also report on four new transiting candidates which we have found in the data. } ",
        "introduction": "The field of extrasolar planets is developing rapidly, producing exciting results at an accelerated pace. The discovery of the first extrasolar ``hot Jupiter'' around the nearby solar-type star 51~Peg using precise radial velocity measurements \\citep{may95} spurred a number of discoveries. Chief among these was the discovery of transits around the nearby solar-type star HD~209458 \\citep{char00,hen00}. The success of the radial velocity studies also boosted extrasolar planetary searches using other techniques such as microlensing and transits. Currently more than 60 transiting extrasolar planets are known\\footnote{See http://exoplanets.eu/}. Many candidates were discovered by the OGLE team who carried out systematic searches, monitoring millions of stars along fields located in the Milky Way disk \\citep{uda02a,uda02b,uda03,pont08}. Of more than 200 transiting candidates, already seven OGLE transits have been confirmed as being due to planets: OGLE-TR-10 \\citep{kon05}, OGLE-TR-56 \\citep{kon03}, OGLE-TR-111 \\citep{pont04}, OGLE-TR-113 \\citep{bou04,kon04}, OGLE-TR-132 \\citep{bou04}, OGLE-TR-182 \\citep{pont08}, OGLE-TR-211 \\citep{uda08}. Most of the other targets are eclipsing low-mass stars or brown dwarfs, or due to blends of normal stars, triplets, etc.  Why such interest in observing transiting candidates? The radial velocities give orbital parameters such as period, semi-major axis, and projected mass ($M~{\\rm sin}~i$). The transits give not only the orbital parameters like period and inclination, but also the planet sizes: the eclipse amplitude is simply $(r/R)^2$. Thus, from combined radial velocities and transits we know the mass of the planet without the inclination ambiguity, and the radius, which gives a density. The few planets so studied appear to be indeed inflated gaseous planets, i.e. ``hot Jupiters''. One difficulty is that the derived planet radii are only good to 10-15\\%. More accurate transit photometry is needed to improve those estimates, as argued by \\cite{mou04}. For example the discoveries by \\cite{pont05a,pont05b} of planet-sized stars around OGLE-TR-106 and OGLE-TR-122, help to constrain the models of planetary systems. These planets are under intense irradiation, which inflates their sizes, depending on their own orbital parameters and intrinsic characteristics \\citep{bar03,burr02}. In some way the OGLE planets constitute the extreme cases, because of their very short periods.  We conducted a programme to monitor photometrically the OGLE transit candidates, and here we present precise photometry for these transits. Some of the objects, namely OGLE-TR-109, OGLE-TR-111, OGLE-TR-113 have been already analyzed by \\cite{fer06}, \\cite{min07}, \\cite{diaz07}, respectively. In this paper we present an analysis of all OGLE transits in the VIMOS images. We also have searched for new transits.  Section~2 gives details on the observations, selection and properties of the sample. Section~3 describes reductions of the data. In sections~4-8 we present the results obtained from the observed light curves of the OGLE transits. Section~9 describes new transiting candidates we have found in the data. Finally, section~10 states our main conclusions. ",
        "conclusions": "$V$-band images from VLT/VIMOS were used to obtain light curves of extrasolar planetary transits: OGLE-TR-111, OGLE-TR-113, and candidate planetary transits: OGLE-TR-109, OGLE-TR-159, OGLE-TR-167, OGLE-TR-170, OGLE-TR-171. With difference imaging photometry we were able to achieve millimagnitude errors in the individual data points. The following seven OGLE transits were recorded as full events: 106, 109, 111, 113, 159, 170, 171. Four transits were detected as partial events: 86, 91, 110, 167. All full and partial transits but OGLE-TR-91 and OGLE-TR-109 were observed at the predicted transit times. No transits were recorded for 19 objects.  Based on the shape of the obtained light curves and some results from spectroscopic follow-up studies we show that the objects OGLE-TR-111 and OGLE-TR-113 are probably the only OGLE stars in the sample to host planets.  In the paper we also report on four new transiting candidates we have found in the VIMOS data. One of the events, transit-4, with the duration time of about 3.7~h and the amplitude of about 0.02~mag, is a particularly good candidate for a planetary transit. Faintness of the object (17.8~mag in $V$) may severely hamper spectroscopic verification. However, all four candidates require photometric follow-up studies to look for their periodic nature first."
    },
    "0910/0910.3683_arXiv.txt": {
        "abstract": "The next generation of proposed galaxy surveys will increase the number of galaxies with photometric redshift identifications by two orders of magnitude, drastically expanding both the redshift range and detection threshold from the current state of the art.  Obtaining spectra for a fair sub-sample of this new data could be cumbersome and expensive.  However, adequate calibration of the true redshift distribution of galaxies is vital to tapping the potential of these surveys to illuminate the processes of galaxy evolution, and to constrain the underlying cosmology and growth of structure.  We examine here a promising alternative to direct spectroscopic follow up: calibration of the redshift distribution of photometric galaxies via cross-correlation with an overlapping spectroscopic survey whose members trace the same density field.  We review the theory, develop a pipeline to implement the method, apply it to mock data from N-body simulations, and examine the properties of this redshift distribution estimator.  We demonstrate that the method is generally effective, but the estimator is weakened by two main factors.  One is that the correlation function of the spectroscopic sample must be measured in many bins along the line of sight, which renders the measurement noisy and interferes with high quality reconstruction of the photometric redshift distribution.  Also, the method is not able to disentangle the photometric redshift distribution from evolution in the bias of the photometric sample.  We establish the impact of these factors using our mock catalogs.  We conclude it may still be necessary to spectroscopically follow up a fair subsample of the photometric survey data.  Nonetheless, it is significant that the method has been successfully implemented on mock data, and with further refinement it may appreciably decrease the number of spectra that will be needed to calibrate future surveys. ",
        "introduction": "\\label{sec:intro} It is essential to calibrate the true redshift distribution of galaxies in a photometric survey if the survey is to be utilized to its full potential.  One application of survey data that requires a detailed understanding of the distribution of galaxies is weak gravitational lensing tomography. The shearing of the shapes of distant galaxies via weak gravitational lensing is a powerful cosmological probe that can be used to study the distribution of dark matter, the nature of dark energy, the formation of large scale structures in the universe, as well as fundamental properties of elementary particles and potential modifications to the general theory of relativity (recent studies include \\cite{2008ARNPS..58...99H}, \\cite{2009MNRAS.395..197T}, \\cite{2009A&A...497..677K}, \\cite{2009PhRvD..79b3520I}). Cosmic shear measurements statistically examine minute distortions in the orientations of high redshift galaxies, whose shapes have been sheared by intervening dark matter structures.   Although weak gravitational lensing provides only an integrated measure of the intervening density field, using source populations at different redshifts permits some degree of three dimensional reconstruction, known as tomography.  The distortions are small (at the 1\\% level) and the intrinsic orientation of the source galaxies is unknown, thus large galaxy samples are required to map the density field and probe the growth of density fluctuations with precision.   Existing cosmic shear measurements have already constrained the amplitude of the dark matter fluctuations at the 10\\% level \\cite{2004MNRAS.347..895H} \\cite{2004ApJ...605...29R} \\cite{2005MNRAS.359.1277M} \\cite{2006A&A...452...51S} and there are many exciting galaxy survey proposals that will increase the available number of source galaxies by two orders of magnitude including DES, DUNE, Euclid, LSST, PanStarrs, SNAP, and Vista.  Because these large galaxy surveys will have photometric rather than spectroscopic redshift identifications, the community has carefully attended to fine tuning the calibration of the photometric redshifts, minimizing biases and catastrophic errors \\cite{2008MNRAS.386..781M}, \\cite{2008MNRAS.390..118L}, \\cite{2009MNRAS.398.2012F}, \\cite{2009AJ....138...95X}, \\cite{2009arXiv0908.4085G}.  Unlike experiments that use the galaxy positions to directly trace the underlying dark matter distribution, such as baryon acoustic oscillation studies, weak lensing analyses do not require a precise redshift identification for each individual source galaxy.  It is sufficient to accurately determine the redshift distribution of the sources. However, lensing measurements are extremely sensitive both to error and bias in the source distribution \\cite{2002PhRvD..65f3001H}.  Attaining an accurate source distribution will be crucial if weak lensing measurements are to be competitive with other cosmological probes in constraining the cosmological parameters.  Another example where calibration of the true distribution of galaxies may be essential is in using the abundances and clustering of different galaxy populations to connect galaxies at late times to their potential progenitors at early times (as in e.g. \\cite{2009ApJ...696..620C}).  Such studies also utilize the luminosity functions of galaxies in each redshift slice, and sometimes divide these into different rest-frame color bins.  To avoid potential systematics in inferences made about galaxy evolution, it will be necessary to know if some fraction of the population in a given photometric redshift bin is actually living at a different redshift, especially if there is an asymmetry in such errors that depends on color.  Recently, an alternative approach to attaining an accurate source redshift distribution has been proposed in \\cite{2008ApJ...684...88N}.   This method is similar to cross-correlation techniques used in \\cite{1985MNRAS.212..657P} and \\cite{2006ApJ...644...54M}, and the idea has also been studied theoretically in \\cite{2006ApJ...651...14S} and \\cite{2009arXiv0902.2782B}.   A similar technique was used in \\cite{2009A&A...493.1197E} to check the redshift distribution for interlopers. Similar in spirit, the analysis of \\cite{2009arXiv0910.2704Q} uses close angular pairs of photometric galaxies to constrain the photometric errors without the use of a spectroscopic sample.  The cross-correlation method determines the photometric redshift distribution by utililizing the cross correlation of the galaxies in the photometric sample with an overlapping spectroscopic sample that traces the same underlying density field. One advantage of this approach is that the spectroscopic sample used to calibrate the photometric redshift distribution can be comprised of bright rare objects such as quasars or Luminous Red Galaxies (LRGs) whose spectra are relatively easy to obtain, and indeed may already exist in legacy data.   Spectra could also be obtained for emission line galaxies (ELGs), which are easy to follow up but may not represent a fair subsample.  Another advantage is that catastrophic redshift errors in the photometry do not systematically bias the redshift distribution estimate, they merely contribute to the noise.  The cross-correlation method makes use of two observables, the line-of-sight projected angular cross-correlation between the photometric and spectroscopic samples $w_{ps}(\\theta)$, and the three dimensional autocorrelation function of the spectroscopic sample $\\xi_{ss}(r)$.   By postulating a simple proportionality between the autocorrelation function of the spectroscopic objects and the three dimensional cross-corelation function between the two samples $\\xi_{ps}(r) \\propto \\xi_{ss}(r)$, it is potentially possible to infer a very accurate redshift distribution for the photometric sample.  This assumption is guaranteed to be valid if the spectroscopic sample is a sub-sample of the photometric population, but may be problematic if the two sets of tracers have different bias functions with respect to the dark matter.  In this paper we develop a pipeline to apply the cross-correlation method.  In section \\ref{sec:theory} we review the theory and explain how the method works.  We highlight its strengths and examine potential drawbacks and systematic effects.  As a proof of concept, in section \\ref{sec:sims} we use the halo model to populate N-body simulations with mock photometric and spectroscopic galaxy data to quantify the properties of this redshift distribution estimator.  We examine the extent to which different bias functions interfere with the reconstruction of the true distribution of photometric galaxies. In section \\ref{sec:disc} we discuss inherent tradeoffs, outline outstanding theoretical questions, and draw our conclusions.  We leave a more detailed discussion of error propagation to the appendix. ",
        "conclusions": "\\label{sec:disc} We have outlined the theory behind the cross-correlation method for calibrating the redshift distribution of objects with photometric redshifts, and developed a pipeline than can be used to apply the method to survey data.  We have created mock simulations to test the pipeline. We have succeeded in reconstructing the redshift distribution of the mock photometric galaxies using the angular cross correlation of these galaxies with an overlapping spectroscopic sample (whose redshifts are known). We have not used any redshift information about the photometric sample.  We have demonstrated the validity of the method. We have also identified the aspects that are likely to be the limiting factors in relying upon this method to provide accurate redshift distribution information.  These limiting factors are  1) that the spectroscopic sample must be binned along the line of sight, causing their correlation functions to be noisy and interfering with the reconstruction, and 2) that the bias evolution of the photometric sample cannot be disentangled from the redshift distribution that is reconstructed, which may force the necessity for the follow up of a fair subsample of galaxies. Improved modeling of theoretical priors could yield large dividends if these factors can be mitigated.  The analysis has revealed a number of trade-offs that exist in application of this method.   For a given set of calibrators, there is a trade-off between the resolution (bin spacing) of the redshift distribution reconstruction, and the error bars on any individual point.  Thus if a population of catastrophic outlier were discovered, for example, it would be interesting to investigate whether it is more important to know how many there are, or at exactly which redshift they lie.  We also find a trade-off between the number of calibrators and the quality of the reconstruction.  As the number of available spectra are decreased, the spectroscopic redshift bins need to widen to maintain equivalent signal strength.  This in turn affects how finely the reconstructed redshift distribution can be sampled.  Bimodal behavior may be lost, and with insufficient priors (such as the smoothness prior we implemented here), false bimodal behavior may appear in the form of anti-correlated adjacent points.  Widening the bins to compensate for fewer spectra also introduces another difficulty.  Recall that the error in the angular cross-correlation function goes as $\\delta w_{ps} \\sim 1/ \\sqrt{n_{\\rm pairs}}$.  This means that \\begin{eqnarray} \\frac{\\delta w_{ps}}{w_{ps}}=\\frac{1}{\\sqrt{n_{\\rm pairs}}} \\left[ \\int_0^\\infty \\xi_{ps} \\phi_p d\\chi\\right]^{-1} \\end{eqnarray} For the purposes of this order of magnitude argument, suppose $\\phi_{p}$ were constant, then to integrate to 1 requires $\\phi_{p}(\\chi) \\propto 1/\\Delta\\chi_{\\rm rb}$.  The subscript rb is used to indicate the width of the reconstruction bin (not the spectroscopic bins). When we remove it from the integral we are left with an expression that is $w_p(r_p)$ \\begin{eqnarray} \\frac{\\delta w_{ps}}{w_{ps}}=\\frac{1}{\\phi_p \\sqrt{n_{\\rm pairs}}} \\left[ \\int_0^\\infty \\xi_{ps} d\\chi\\right]^{-1} \\nonumber \\\\ \\sim \\frac{\\Delta\\chi_{\\rm rb}}{\\phi_p w_p(r_p) \\sqrt{n_{\\rm pairs}}} \\end{eqnarray} If the reconstruction bin is widened by a factor of two, the number density of photometric galaxies will have to be increased by a factor of 4 to preserve the same accuracy in the cross correlation measurement.   Therefore in this method there is also a trade-off between the number of spectra that the survey can afford versus the number of photometric galaxies they can afford.  In this analysis we have not made any use of the photometric redshifts, which although not sufficiently accurate to determine the true redshift distribution of the population, still contain significant information about it.  Modern techniques such as in \\cite{2009arXiv0908.4085G} have made it possible to assign a probability distribution for the redshift of each individual galaxy in a photometric survey, rather than a single best estimate and error bar.   We propose that combining these probabilities for all the galaxies in a given redshift bin constitutes a reasonably powerful prior, that can take the place of the smoothing we have introduced in this analysis.  This is fortunate because the solution is frequently wrecked by the smoothing criterion, although the analysis cannot be performed without it.  Another detail that we leave for a future study is that the spectroscopic sample need not be comprised of a single rare population, it is quite conceivable to target certain regions of redshift space that are known to be problematic more heavily.  It must also be possible to use less rare tracers in regions with smaller volume.  We leave such optimization to the future.  There are a few important factors that we have neglected in this analysis.  As pointed out by \\cite{2009arXiv0902.2782B}, weak gravitational lensing will induce correlations between the positions of calibrator galaxies in the foreground with photometric galaxies in the background, and vice versa.  This will need to be carefully controlled for the method to reliably calibrate redshift distributions.  We also have not mentioned the complication of the integral constraint, which as discussed in e.g. \\cite{2007APh....26..351H} can lead to very significant errors when the volume and the scales in the correlation function are of comparable size.  This may be as significant an issue as evolution in the large scale bias of the photometric population, though it may be mitigated if integral constraint errors are correlated between photometric and spectroscopic samples.  Fortunately, improved estimators exist (in \\cite{2007MNRAS.376.1702P} for example) and should certainly be incorporated into the pipeline.  Having now demonstrated that the method is viable, with further refinement the cross-correlation method applied in conjunction with direct follow up surveys may significantly reduce the number of spectra that are required to calibrate photometric redshift distributions to the desired accuracy.  There are a few other advantages that we have not touched upon in detail.  As we have shown, the cross correlation method can be used to calibrate the redshift distribution all the way along the line of sight, and as such is uniquely suited to detection and calibration of catastrophic redshift errors, even if these errors are so rare that they are missed by conventional follow up.  Also, the reconstruction should not be adversely affected by redshift deserts, regions where no spectroscopic redshifts are available.  We are optimistic that this technique will provide a useful complementary approach to conventional calibration techniques, and every effort should be made to refine the method further."
    },
    "0910/0910.3010_arXiv.txt": {
        "abstract": "Measurement of the spatial distribution of neutral hydrogen via the redshifted 21~cm line promises to revolutionize our knowledge of the epoch of reionization and the first galaxies, and may provide a powerful new tool for observational cosmology from redshifts $1<z<4$. In this review we discuss recent advances in our theoretical understanding of the epoch of reionization (EoR), the application of 21~cm tomography to cosmology and measurements of the dark energy equation of state after reionization, and the instrumentation and observational techniques shared by 21~cm EoR and post reionization cosmology machines. We place particular emphasis on the expected signal and observational capabilities of first generation 21~cm fluctuation instruments. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.4406_arXiv.txt": {
        "abstract": "We  present a  catalogue of  1602 white  dwarf-main  sequence binaries (WDMS) from the spectroscopic Sloan  Digital Sky Survey Data Release 6 (SDSS DR6).   Among these  we identify 440  as new WDMS  binaries.  We select WDMS  binary candidates from template fitting  all 1.27 million DR6  spectra,  using  combined  constraints  in  both  $\\chi^{2}$  and signal-to-noise ratio.  In addition,  we use Galaxy Evolution Explorer (GALEX) and  UKIRT Infrared Sky  Survey (UKIDSS) magnitudes  to search for  objects in which  one of  the two  components dominates  the SDSS spectrum.   We use  a decomposition/fitting  technique to  measure the effective temperatures, surface gravities, masses and distances to the white  dwarfs, as  well as  the spectral  types and  distances  to the companions in our catalogue.   Distributions and density maps obtained from these stellar parameters are  then used to study both the general properties  and the  selection effects  of WDMS  binaries in  SDSS.  A comparison between the distances measured  to the white dwarfs and the main  sequence  companions  shows  $d_\\mathrm{sec}>d_\\mathrm{wd}$  for $\\sim$1/5 of  the systems,  a tendency found  already in  our previous work. The hypothesis that  magnetic activity raises the temperature of the inter-spot  regions in  active stars that  are heavily  covered by cool spots,  leading to  a bluer optical  colour compared  to inactive stars, remains the best explanation  for this behaviour.  We also make use of SDSS-GALEX-UKIDSS magnitudes to investigate the distribution of WDMS binaries, as well as their white dwarf effective temperatures and companion  star  spectral types,  in  ultraviolet  to infrared  colour space.  We show  that WDMS binaries can be  very efficiently separated from single main sequence stars and white dwarfs when using a combined ultraviolet, optical, and infrared colour selection.  Finally, we also provide  radial   velocities  for  1068  systems   measured  from  the \\Lines{Na}{I}{8183.27,8194.81} absorption doublet and/or the H$\\alpha$ emission line.  Among the  systems with multiple SDSS spectroscopy, we find   five  new  systems   exhibiting  significant   radial  velocity variations,   identifying    them   as   post-common-envelope   binary candidates. ",
        "introduction": "\\label{s-intro}  Binaries  containing  a  white  dwarf  primary plus  a  main  sequence companion  were initially  main sequence  binaries in  which  the more massive  star evolved  through  the  giant phase  and  became a  white dwarf. In  the majority  of cases the  initial separation of  the main sequence binary is wide enough to allow the evolution of both stars as if they were single.  A small  fraction is believed to undergo a phase of dynamically unstable mass transfer once the more massive star is on the giant  branch or  the asymptotic giant  branch \\citep{webbink84-1, dekool92-1, willems+kolb04-1}.  As a consequence of this mass-transfer the envelope of the giant will engulf its core and the companion star, i.e.   the  system is  entering  a  common  envelope phase  \\citep[CE, e.g.][]{livio+soker88-1,      iben+livio93-1,      taam+sandquist00-1, webbink07-1}.  Friction  inside this envelope causes  a rapid decrease of  the  binary separation.   Henceforth  orbital  energy and  angular momentum are extracted from the  binary orbit and lead to the ejection of the  envelope, exposing a post-common-envelope binary  (PCEB). As a consequence  the  orbital  period  distribution of  WDMS  binaries  is clearly bi-modal,  with PCEBs  concentrated at short  orbital periods, and   wide  WDMS   binaries  (non-PCEBs)   at  long   orbital  periods \\citep{willems+kolb04-1}.  After  the ejection of  the envelope, close WDMS  binaries  evolve  to  shorter orbital  periods  through  angular momentum  loss via  magnetic braking  and/or through  the  emission of gravitational  waves.  WDMS  binaries  include progenitors  of a  wide range  of  astronomical  objects  such as  cataclysmic  variables  and super-soft X-ray  sources, with some of those  objects likely evolving at later stages into type Ia supernova.  Population  synthesis models  have  been developed  for  a variety  of binary   stars   undergoing   CE   evolution   \\citep{dewi+tauris00-1, nelemans+tout05-1,      politano+weiler06-1,      politano+weiler07-1, davisetal08-1, davisetal09-1}.  However, the theoretical understanding of  both  CE  evolution  and  magnetic  braking  is  currently  poorly constrained   by  observations   \\citep{schreiber+gaensicke03-1},  and progress on this front is most  likely to arise from the analysis of a large  sample of  PCEBs that  are  well-understood in  terms of  their stellar  components.   WDMS binaries  appear  most  promising in  that respect, as  their stellar components  are relatively simple,  and the SDSS \\citep{yorketal00-1, stoughtonetal02-1, abazajianetal09-1} offers the possibility  to dramatically increase the number  of WDMS binaries available for detailed follow-up studies.  Up to date there exist four compilations of SDSS  WDMS binaries, namely \\citet{eisensteinetal06-1} (obtained   as    part   of    their   search   of    white   dwarfs), \\cite{silvestrietal07-1}          (which         contains         also \\citet{raymondetal03-1,silvestrietal06-1} as a subset, and was claimed to be complete for the SDSS DR5), \\citet{augusteijnetal08-1} (obtained from   SDSS  DR5  using   colour  cuts   plus  proper   motions),  and \\citet{helleretal09-1} (obtained  from a search  for sdB stars  in the SDSS DR6).  Here we initiate  a comprehensive study to compile  a master sample of spectroscopic WDMS  binaries from SDSS.  In this  first publication we make      use       of      the      SDSS       spectroscopic      DR6 \\citep{adelman-mccarthyetal08-1}  to create a  catalogue of  1602 WDMS binaries and candidates that were serendipitously observed.  This list is  not  only  a superset  of  all  those  previous SDSS  WDMS  binary compilations, but also includes 440 new WDMS binaries.  In addition we provide a coherent  analysis of the system parameters  of both stellar components, as well  as an extension to GALEX/UKIDSS  colours to study the properties  of SDSS WDMS binaries.   In a parallel  effort, and as part  of SEGUE  (the  SDSS Extension  for  Galactic Understanding  and Exploration), some of  us carried out a dedicated  program to identify $\\sim300$    WDMS    binaries    containing    cold    white    dwarfs \\citep{schreiberetal07-1},  a population  clearly  underrepresented in previous samples of  WDMS binaries.  A detailed analysis  of the SEGUE population  of  WDMS  is  in  preparation (Schreiber  et  al.   2009). Finally,  we plan  to summarise  our  efforts by  presenting the  WDMS binary     content    of     the    final     SDSS     data    release \\citep[DR7,][]{abazajianetal09-1}. This last paper will be accompanied by a public online data base of WDMS binaries.  The final  master sample of  ``all'' spectroscopic SDSS  WDMS binaries will  form a  superb database  for  future follow-up  studies of  WDMS binaries.    In   particular,   analysing   the   fraction   and   the characteristics  of the  PCEBs  among the  WDMS  binaries may  provide strong  constraints on  current theories  of compact  binary evolution \\citep{rebassa-mansergasetal07-1,                    schreiberetal08-1, rebassa-mansergasetal08-1,        pyrzasetal09-1,       nebotetal09-1, schwopeetal09-1}.  The structure of the paper  is as follows.  In Sect.\\,\\ref{s-ident} we present  our method  of identifying  WDMS binaries,  and  estimate the completeness  of   the  sample  in   Sect.\\,\\ref{s-completeness}.   In Sect.\\,\\ref{s-final} we provide our final catalogue.  Using a spectral decomposition/model  atmosphere   analysis,  we  derive   white  dwarf effective  temperatures,  surface  gravities, masses,  companion  star spectral  types,  and distances  in  Sect.\\,\\ref{s-param}. We  compare these  stellar  parameters  to  those  obtained in  other  studies  in Sect.\\,\\ref{s-comparison},  and study  the selection  effects  of WDMS binaries  within  SDSS  in  Sect.\\,\\ref{s-effects}.  We  provide  then colour-colour  diagrams and  colour-colour cuts  for WDMS  binaries in Sect.\\,\\ref{s-colors}   and  Sect.\\,\\ref{s-cuts},   respectively.   We finally measure  the \\Lines{Na}{I}{8183.27,8194.81} absorption doublet and/or the H$\\alpha$  emission radial velocities in Sec.\\,\\ref{s-rvs}, and summarise our results in Sect.\\,\\ref{s-conc}.  \\begin{figure*} \\begin{center} \\includegraphics[angle=-90,width=\\textwidth]{wdms.ps} \\caption{Six examples  of previously known WDMS binaries  used in this work as WDMS binary  templates. SDSS names and MJD-PLT-FIB identifiers are indicated for each of them. White dwarf effective temperatures and spectral  type of  the companions  (see Sec.\\,\\ref{s-param})  are also indicated in each panel.} \\label{f-templ} \\end{center} \\end{figure*} ",
        "conclusions": "\\label{s-conc}  We  have  presented  a  catalogue  of  1602  WDMS  binaries  from  the spectroscopic  SDSS   DR6.   We  have   used  a  decomposition/fitting technique  to measure the  effective temperatures,  surface gravities, masses  and distances to  the white  dwarfs, as  well as  the spectral types and distances to the companions in our catalogue.  Distributions and density maps obtained from these stellar parameters have been used to study both the general properties and the selection effects of WDMS binaries in SDSS.  A comparison  between the distances measured to the white    dwarfs    and   the    main    sequence   companions    shows $d_\\mathrm{sec}>d_\\mathrm{wd}$ for $\\sim$1/5  of the systems.  We also have  made use  of  GALEX, SDSS  and  UKIDSS magnitudes  to study  the distribution  of  WDMS binaries  in  colour-colour  space and  present simple colour-cuts  that allow to clearly separate  WDMS binaries from other stellar objects. Finally, we have measured radial velocities for 1068  WDMS binaries  measured from  the \\Lines{Na}{I}{8183.27,8194.81} absorption  doublet and/or  the  H$\\alpha$ emission  line.  Among  the systems  with  multiple  SDSS  spectroscopy,  we find  five  new  WDMS binaries  showing significant  radial velocity  variations identifying them as  PCEB candidates.  The  here presented new, updated,  and most complete catalogue of WDMS binaries  from the SDSS represents a superb data  base  for  future   follow-up  studies  that  may  significantly contribute  to a  better understanding  of close  compact  binary star evolution."
    },
    "0910/0910.1223_arXiv.txt": {
        "abstract": "We report the results of a search for 12.2-GHz methanol maser emission, targeted towards 113 known 6.7-GHz methanol masers associated with 1.2-mm dust continuum emission. Observations were carried out with the Australia Telescope National Facility (ATNF) Parkes 64-m radio telescope in the period 2008 June 20 - 25. We detect 68 12.2-GHz methanol masers with flux densities in excess of our 5-$\\sigma$ detection limit of 0.55~Jy, 30 of which are new discoveries. This equates to a detection rate of 60 per cent, similar to previous searches of comparable sensitivity.  We have made a statistical investigation of the properties of the 1.2-mm dust clumps with and without associated 6.7-GHz methanol maser and find that 6.7-GHz methanol masers are associated with 1.2-mm dust clumps with high flux densities, masses and radii. We additionally find that 6.7-GHz methanol masers with higher peak luminosities are associated with less dense 1.2-mm dust clumps than those 6.7-GHz methanol masers with lower luminosities. We suggest that this indicates that more luminous 6.7-GHz methanol masers are generally associated with a later evolutionary phase of massive star formation than less luminous 6.7-GHz methanol maser sources.  Analysis of the 6.7-GHz associated 1.2-mm dust clumps with and without associated 12.2-GHz methanol maser emission shows that clumps associated with both class~II methanol maser transitions are less dense than those with no associated 12.2-GHz methanol maser emission. Furthermore, 12.2-GHz methanol masers are preferentially detected towards 6.7-GHz methanol masers with associated OH masers, suggesting that 12.2-GHz methanol masers are associated with a later evolutionary phase of massive star formation.  We have compared the colours of the GLIMPSE point sources associated with the maser sources in the following two subgroups: 6.7-GHz methanol masers with and without associated 12.2-GHz methanol masers; and 6.7-GHz methanol masers with high and and those with low peak luminosities. There is little difference in the nature of the associated GLIMPSE point sources in any of these subgroups, and we propose that the masers themselves are probably much more sensitive than mid-infrared data to evolutionary changes in the massive star formation regions that they are associated with.  We present an evolutionary sequence for masers in high-mass star formation regions, placing quantitative estimates on the relative lifetimes for the first time. ",
        "introduction": "Interstellar masers are one of the most readily observed signposts of star formation, particularly emission from the hydroxyl (OH), water (H$_2$O) and methanol (CH$_3$OH) molecules as they are relatively common and intense. Because these masers arise at radio frequencies, they are not affected by the dense optical-obscuring gas and dust present at the early stages of massive star formation thereby allowing us to probe these stars at the earliest stages.  The masers are powerful probes of the kinematics of the star formation regions, but we are only recently beginning to be able to use the masers to investigate aspects such as the physical conditions through comparison of observational data and maser pumping theories \\citep{Cragg01,Sutton01}. The different maser species favour different physical conditions and are thought to trace different evolutionary phases of massive star formation \\citep{Ellingsen07}.  Since many sources show emission from multiple maser transitions or species there must be significant overlap for the evolutionary phase traced by the most common types of masers.  Interstellar masers have been detected towards more than a thousand star formation regions in our Galaxy from transitions of methanol, hydroxyl and water. Transitions of class II methanol masers have some advantages over water and OH masers as they exclusively trace sites of massive star formation \\citep{Minier03,Xu08} while water and OH are known to also be associated with other astrophysical objects such as evolved stars and supernova remnants. Two of the strongest maser species found towards star formation regions are transitions of methanol at 6.7- and 12.2-GHz, both of which are known to trace an early evolutionary stage of massive star formation \\citep{Ellingsen06}. Complementary observations of these two strong methanol maser transitions are especially useful in probing the physical conditions of the environments in which they arise, as they are known to typically be co-spatial \\citep{Menten92,Norris93}.  To date there have been several searches for 12.2-GHz methanol maser emission towards 6.7-GHz methanol maser emission \\citep*[e.g.][]{Caswell95b,Blas04}. \\citet{Caswell95b} targeted their search for 12.2-GHz methanol masers towards 238 6.7-GHz methanol masers that had been detected towards sites of OH masers and achieved a detection rate of 55 per cent. \\citet{Blas04} directed their observations at 12.2-GHz towards 6.7-GHz methanol masers that were observed towards a sample of {\\em IRAS} (Infrared Astronomy Satellite) selected sources \\citep{Szy00} as well as 6.7~GHz methanol masers that were detected in a blind search that had relatively poor positional accuracy and sensitivity \\citep{Szy02}.  Searches such as those by \\citet{Caswell95b} and \\citet{Blas04} have shown that 12.2-GHz methanol maser emission is only rarely brighter than the associated 6.7-GHz methanol maser emission. This coupled with the fact that there have been no serendipitous detections of 12.2-GHz masers without a 6.7-GHz counterpart mean that it is unlikely that an unbiased Galactic survey would uncover many, if any, more 12.2-GHz methanol masers than would be yielded from a search targeted towards 6.7-GHz methanol maser sources.  Recently, a small subset of the strong 6.7-GHz methanol masers from the sources observed by \\citet{Hill05} at 1.2-mm, were searched for the presence of 12.2-GHz methanol masers with the University of Tasmania 26~m Mount Pleasant radio telescope \\citep{Lewis07}. \\citet{Lewis07} detected 17 12.2-GHz methanol masers towards the 27 sites searched with a detection limit of $\\sim$4~Jy. \\citet{Hill05} observed 131 regions that were suspected of undergoing massive star formation, using the presence of previously identified 6.7-GHz methanol masers and/or an ultracompact \\ionhy region to select these regions. A total of 404 1.2-mm dust clumps were identified, with 113 of these associated with 6.7-GHz methanol masers, 35 of which are associated with radio continuum emission \\citep{Walsh98}.   Statistical analysis by \\citet{Lewis07} of the properties of 1.2-mm dust clumps associated with 6.7-GHz methanol maser emission with and without 12.2-GHz methanol masers showed that those 1.2-mm dust clumps with both 6.7-GHz and 12.2-GHz were less massive and less dense than the 1.2-mm clumps that were associated with only 6.7-GHz methanol maser emission.  Comparison between the 6.7-GHz methanol masers with detected 12.2-GHz with those with no detectable 12.2-GHz maser emission, showed that 12.2-GHz methanol masers were preferentially found at those 6.7-GHz methanol maser sites with associated OH emission (to a better than 95 per cent confidence). As OH masers are expected to trace a generally later stage of star formation \\citep{FC89,Cas97} than 6.7-GHz methanol masers. As the sample size of the \\citet{Lewis07} observations was small, was subject to a 6.7-GHz flux density bias and had relatively poor sensitivity, the aim of the current work was to confirm these results by testing a large sample with a much lower 12.2-GHz detection limit.   Here we present a targeted search for 12.2-GHz methanol masers, with a greatly increased sample size and sensitivity (c.f. \\citet{Lewis07}), towards 113 known 6.7-GHz methanol masers which have been previously targeted for 1.2-mm dust continuum emission \\citep{Hill05}. Statistical analysis of the 1.2-mm dust clumps devoid of either methanol maser transition and those associated with 12.2-GHz and/or 6.7-GHz methanol maser sources has been carried out, with a particular emphasis on the investigation of the relative evolutionary stage each class of source is tracing. ",
        "conclusions": "We have searched 113 known 6.7-GHz methanol masers for associated 12.2-GHz methanol maser emission with the ATNF Parkes 64-m radio telescope. We have detected 12.2-GHz methanol masers towards 68 of these sources; a detection rate of 60 percent. We estimate the life-time of 12.2-GHz methanol maser emission to be of the order of 2.1 $\\times$ 10$^{4}$ years. We find that for the majority of sources, the peak velocity of the 12.2-GHz methanol maser emission is the same as the 6.7-GHz methanol maser and, for the remainder, the 12.2-GHz peak emission is coincident with a secondary feature present in the 6.7-GHz methanol maser spectrum. Comparison between the peak flux densities of the respective methanol maser transitions has shown that median ratio of 6.7- to 12.2-GHz peak flux density is 1:5.9, similar to previous searches. We additionally find that the ratio of 6.7- to 12.2-GHz methanol maser peak flux density has some dependence on 6.7-GHz methanol maser luminosity, with the more luminous 6.7-GHz sources having comparatively weaker 12.2-GHz methanol maser emission and vice versa.  All of the target sources have previously been observed at 1.2-mm for dust continuum emission \\citep{Hill05}. We have carried out statistical analysis of the dust clump properties of sources in the following four subgroups compared the 1.2-mm dust clump properties: dust clumps with and without associated 6.7-GHz methanol maser emission; dust clumps associated with highly luminous 6.7-GHz methanol masers compared with those of low luminosity; dust clumps associated with 6.7-GHz methanol maser emission with and without associated 12.2-GHz methanol maser emission; and dust clumps with and without associated radio continuum. Our statistics show that there are many differences between the sources that fall into these subgroups; consistent with the different source combinations and characteristics being associated with different stages of massive star formation. An evolutionary scenario for the common maser species associated with high-mass star formation regions is presented.  The positions of the 6.7-GHz methanol masers have been compared to sources in the GLIMPSE point source catalogue. We find that 57 per cent of the 6.7-GHz methanol masers that lie within the GLIMPSE regions are associated with a GLIMPSE point source. Comparison between the GLIMPSE colours associated with sources that exhibit emission at both 6.7- and 12.2-GHz and those exclusively at 6.7-GHz has shown that there is little difference between the colours of the associated point sources.  We will present detailed investigation of the relative intensities and velocity ranges of 12.2-GHz methanol masers detected towards the 6.7-GHz methanol masers detected in the southern hemisphere component of the MMB survey \\citep{Green08} in a subsequent paper. The near simultaneous observations of hundreds of 6.7- and 12.2-GHz methanol masers collected from the MMB survey and follow-up observations will be used to make a detailed investigation of the range and distribution of the intensity ratio for these two transitions.  Combining this with our evolutionary studies we will obtain unique insight into the changes in physical conditions responsible for the presence/absence of the different maser methanol transitions.  Our results relating to the evolution of star formation regions will be independently tested in the future through comparison between the unbiased sample of 6.7-GHz methanol masers observed in the MMB survey \\citep{Green08}, followup observations at 12.2-GHz and 870~$\\mu$m dust sources observed as part of ATLASGAL (APEX Telescope Large Area Survey of the Galaxy)\\footnote{http://www.mpifr-bonn.mpg.de/div/atlasgal/}."
    },
    "0910/0910.1976_arXiv.txt": {
        "abstract": "{ We study the evolution of   curvature perturbations and the cosmic microwave background (CMB) power spectrum in the presence of an hypothesized extra anisotropic stress which might arise, for example, from the dark radiation term in brane-world cosmology. We evolve the scalar modes of  such perturbations before and after neutrino decoupling and analyze their effects on the CMB spectrum.  A novel result of this work is that  the cancellation of the neutrino and extra anisotropic stress could lead to a spectrum of residual  curvature perturbations which  is similar to  the observed CMB power spectrum. This implies a possible additional consideration in the determination of cosmological parameters from the CMB analysis. } \\received{\\today} \t\t% \\accepted{\\today} ",
        "introduction": "Anisotropic stress is the traceless component in the energy-momentum tensor and it accounts for  anisotropic momentum flow in the universe. Although  some previous works  have considered  neutrino anisotropic stress and the anisotropic stress of the primordial magnetic field \\cite{1995ApJ...455....7M,2008PhRvD..77f3003G,2006CQGra..23.4991G,2006PhRvD..74f3002G,2008arXiv0806.2018K,2008nuco.confE.226K}, they mainly considered the post neutrino decoupling era. In this work we deduce the evolution of the neutrino anisotropic stress and curvature perturbations both before and after neutrino decoupling which can be important.   In the standard cosmological model, the anisotropic stress is absent before the epoch of neutrino decoupling because rapid interactions among elementary particles dissipate it. After the neutrinos have  decoupled from the cosmic expansion, however, the neutrinos become relativistic free particles and neutrino anisotropic stress can grow gradually. The neutrino anisotropic stress plays an important role in the formation of large scale structure and the CMB fluctuation spectrum.  In this paper, however,  we consider other possible sources for anisotropic stress besides neutrinos or a magnetic field which can be present before neutrino decoupling. In particular,  such terms arise naturally in cosmological theories in higher dimensions.  We follow the evolution of the neutrino anisotropic stress and curvature perturbations both before and after neutrino decoupling in the presence of such anisotropic stress terms and show that curvature perturbations can stay constant on superhorizon scales as does the standard adiabatic mode. However, unlike the standard case, significant evolution of curvature perturbations occurs before neutrino decoupling. We show that, for the right conditions, such curvature perturbations might even reproduce the observed CMB power spectrum implying a possible additional consideration in the extraction of cosmological parameters from the CMB analysis.    One possible source of extra anisotropic stress  is the so-called \"dark radiation\" term in brane-world cosmology.  This term can affect strongly the CMB anisotropies \\cite{2001PhRvD..63h4009L,2002PhRvD..65j4012L,2002PhLB..532..153B,2002PhRvD..66d3521I, 2003PhRvD..68h3518I, 2006PhRvD..73f3527U}. In such brane-world cosmology, there are  two correction terms in the 4D Einstein equation. One is the extrinsic curvature which introduces a quadratic term in the  energy momentum tensor, but is negligible in the  low energy limit. The other  is from the projected Weyl tensor. The application of the five-dimensional conservation condition to the Weyl tensor requires that the energy density term varies with scale factor as $a^{-4}$  to an observer on the brane \\cite{2001PhRvD..63h4009L}.  Hence, it is called the  dark radiation term and remains significant throughout the radiation-dominated epoch.  There is, however,  no intrinsic brane equation to determine the anisotropic stress \\cite{2001PhRvD..63h4009L}. Although there  have been attempts (e.g. \\cite{2003PhRvL..91v1301K}) to constuct simple models for the Weyl anisotropic stress, the solution strongly depends upon unknown physics in the bulk dimension. Hence, in what follows we take a  phenomenological approach and simply adopt the plausible assumption that in the same way that inflation generates scale-free fluctuations characterized by a spectral index $P(k) \\sim k^n$ (with $n \\sim 1$), we can expect that the Weyl anisotropic stress could be characterized by a spectral index  to be determined (constrained) by a fit to the observed CMB power spectrum. Moreover, this dark-radiation term is assumed to scale as $\\propto a^{-4}$ similar to the energy density and pressure as was adopted in Ref. \\cite{2002PhLB..532..153B}.  Another more familiar  example of extra  anisotropic stress is that of  a primordial magnetic field (PMF). The amplitude of the energy density $B^2/8\\pi$ and magnetic anisotropic stress $\\rho_\\gamma\\pi_B$ of the PMF again both scale as radiation density $\\propto a^{-4}$.  Moreover, such a PMF should be characterized by an amplitude and spectral index, However, the contribution of a PMF to the CMB power spectrum is constrained to be small   because it has a measurable effect on the observed BB mode of the CMB on small angular scales \\cite{2008PhRvD..77d3005Y, 2008PhRvD..78l3001Y}. As shown in Ref.~\\cite{2004PhRvD..70d3011L}, the tensor mode of the primordial magnetic field has two kinds of fluctuations. One is a compensated mode, which arises from the compensation of anisotropic stress, and the other is a passive mode which was generated by the anisotropic stress of the magnetic field before the epoch of neutrino decoupling. The passive tensor mode makes a large contribution to the CMB. Although Ref.~\\cite{2004PhRvD..70d3011L} also studied the vector mode, there was no analysis of the scalar mode in that work. The spectral index of a PMF is also constrained from the effects of the associated gravity waves on  primordial nucleosynthesis \\cite{2000PhRvD..61d3001D}. Hence, a PMF cannot source the scale-free extra anisotropic stress of interest to this work. We only mention it as an example of a power-law anisotropic stress which also scales as $\\propto a^{-4}$.  Here we show that a scale-free  anisotropic stress, if present, could imply an additional consideration in the determination of cosmological parameters from the CMB. As an extreme example, in this paper we show that an extra anisotropic stress could even  lead to a CMB spectrum that agrees with observation. Such a source is a bit contrived as it must be scale invariant as naturally occurs in the inflationary scenario. Thus some kind of mechanism would have to occur to produce the desired scale-invariant anisotropic stress. Our main point is that the possibility exists that such an anisotropic stress  term might be present in the early universe and that it would lead to curvature perturbations which could reproduce the  observed CMB power spectrum. This possibility has not previously been pointed out to our knowledge.  A second point which we make here is that this extra source of anisotropic stress may produce non-Gaussian fluctuations depending  upon the nature of the source of anisotropic stress. The WMAP-5yr analysis, has indicated that there is at least a possibility for  non-Gaussianity although the uncertainty is large \\cite{2008arXiv0803.0547K}. The Planck mission should constrain non-Gaussianity with high precision. If non-Gaussianity were actually observed, it could suggest (e.g.~\\cite{2004PhR...402..103B}) the need for a new cosmological paradigm which allows for  non-Gaussianity such as that described here. At the very least, observational limits on non-Gaussianity could place limits on the hypothesis proposed here, depending upon the source of the extra  anisotropic stress. ",
        "conclusions": "We have found a simple solution for Eq. (\\ref{eta2}) and have shown that if there exists an extra anisotropic stress which scales as $\\rho_\\gamma\\pi_{\\rm ex}\\propto a^{-4}$, the anisotropic stress from neutrinos $\\pi_\\nu$ exactly cancels $\\pi_{\\rm ex}$. Before neutrino decoupling, the curvature perturbations grow logarithmically. After neutrino decoupling, they become constant on superhorizon scales just like the standard adiabatic mode of inflation. This is  because the total anisotropic stress vanishes via a  cancellation. Thus, the resultant CMB spectrum is a superposition of the primary  and passive modes.  As an illustration we have considered the possibility that the passive scalar mode has a scale invariant spectrum.  In this case  the extra anisotropic stress might even produce a power spectrum similar to the observed CMB. This suggests a possible additional consideration in the determination of cosmological parameters from the CMB. Also we note that the  Gaussianity of the CMB fluctuations depends upon the  source for the extra anisotropic stress. Hence, should this extra anisotropic stress be present in the observed power spectrum, it might be detectable by non-Gaussianity in the fluctuations. Future CMB observations of  non-Gaussianity may, therefore,  help to constrain this possibility.  To summarize, our purpose here has been to suggest that such a contribution to the observed spectrum may exist. Thus, further studies on the amplitude and spectrum of the extra anisotropic stress in brane-world cosmology are warranted."
    },
    "0910/0910.2144_arXiv.txt": {
        "abstract": "A new promising development in astroparticle physics is to measure the radio emission from extensive air showers. The particles in the cascade emit synchrotron radiation (30 - 90 MHz) which is detected with arrays of dipole antennas.  Recent experimental efforts are discussed. ",
        "introduction": "An intense branch of astroparticle physics is the study of high-energy cosmic rays to reveal their origin, as well as their acceleration and propagation mechanisms \\cite{behreview,cospar06,wuerzburg}. At energies exceeding $10^{14}$~eV cosmic rays are usually studied by indirect measurements --- the investigation of extensive air showers initiated by cosmic particles in the atmosphere. Different techniques are applied, like the measurements of particle densities and energies at ground level, or the observation of \\Cerenkov and fluorescence light. An alternative technique has been recently revitalized --- the detection of radio emission (at tens of MHz) from extensive air showers at energies exceeding $10^{16}$~eV.  Radio emission from air showers was experimentally discovered in 1965 at a frequency of 44 MHz \\cite{jelleynature}. The early activities in the 1960s and 1970s are summarized in \\rref{allanrev}.  Only recently, fast analog-to-digital converters and modern computer technology made a clear detection of radio emission from air showers possible.  LOPES, a LOFAR Prototype Station had shown that radio emission from air showers can be detected even in an environment with relatively strong radio frequency interference (RFI) \\cite{radionature}. Further investigations of the radio emission followed with LOPES \\cite{badearadio,petrovicinclined,buitinkthunder,niglfreq,nigldirect} and the CODALEMA experiment \\cite{codalema,Ardouin:2006nb}, paving the way for this new detection technique \\cite{Falcke:2008qk}.  Most likely, the dominant emission mechanism of the radio waves in the atmosphere is synchrotron radiation due to the deflection of charged particles in the Earth magnetic field (geosynchrotron radiation) \\cite{huege2003,huege2005,huegefalcke}.  In the frequency range of interest ($30-90$~MHz) the wavelength of the radiation is large compared to the size of the emission region: the typical thickness of the shower disc is about 1 to 2~m only. Thus, coherent emission is expected, which yields relatively strong signals (of the order of a few $\\mu$V/m MHz) at ground level.  Goal of the present activities is to further push the development of radio detection to become a new, independent way to measure the properties of air showers. Ultimate goal is to derive information about the primary, shower-inducing particle from the measurements, such as the energy and mass of the particle as well as the particle direction and point of incidence.  An advantage of radio detection, e.g.\\ with respect to the fluorescence technique, is that showers can be observed with almost 100\\% duty cycle.  The {\\sl arrival direction} can be inferred from the arrival time of the wavefront at the antennas. The measured radio signal at a certain distance from the shower axis is a good estimator for the number of particles at shower maximum, being almost independent of the primary particle type \\cite{huege2008}. This can be used to determine the shower {\\sl energy}. The curvature of the shower front has been investigated \\cite{lafebrecurv,icrc09-lafebre}. It could be shown that the radius of curvature measured at ground level is related to the distance to the shower maximum.  The depth of the shower maximum in the atmosphere \\Xmax is an important observable to determine the mass of the primary particle.  The study indicates a resolution of \\Xmax from radio observations of order of $30-40$~\\gcm2, i.e.\\ comparable to the accuracy of present fluorescence detectors. ",
        "conclusions": ""
    },
    "0910/0910.0836_arXiv.txt": {
        "abstract": "We report on the results of a \\emph{Chandra} search for evidence of triggered nuclear activity within the Cl0023+0423 four-way group merger at $z\\sim0.84$.  The system consists of four interacting galaxy groups in the early stages of hierarchical cluster formation and as such, provides a unique look at the level of processing and evolution already underway in the group environment prior to cluster assembly.  We present the number counts of X-ray point sources detected in a field covering the entire Cl0023 structure, as well as a cross-correlation of these sources with our extensive spectroscopic database.   Both the redshift distribution and cumulative number counts of X-ray sources reveal little evidence to suggest the system contains X-ray luminous AGN in excess to what is observed in the field population.   If preprocessing is underway in the Cl0023 system, our observations suggest that powerful nuclear activity is not the predominant mechanism quenching star formation and driving the evolution of Cl0023 galaxies.  We speculate that this is due to a lack of sufficiently massive nuclear black holes required to power such activity, as previous observations have found a high late-type fraction among the Cl0023 population.  It may be that disruptive AGN-driven outflows only become an important factor in the preprocessing of galaxy populations during a later stage in the evolution of such groups and structures, when sufficiently massive galaxies (and central black holes) have built up, but prior to hydrodynamical processes stripping them of their gas reservoirs. ",
        "introduction": "There is now substantial evidence that environments of intermediate density, such as galaxy groups, play an important role in the transformation of field galaxies into the passively evolving populations found in galaxy clusters. Several studies have found that group populations already exhibit reduced star formation rates (SFR; Lewis et al.~2002; Gomez et al.~2003) and high early-type fractions similar to those observed in denser environments (Zabludoff \\& Mulchaey 1998; Jeltema et al.~2007).  While the physical mechanisms responsible for the preprocessing of galaxies in the group regime are still heavily debated, several recent studies have reported an overdensity of X-ray luminous Active Galactic Nuclei (AGN) on the outskirts of clusters and within the substructure surrounding unrelaxed systems (D'Elia et al.~2004; Cappelluti et al. 2005; Kocevski et al.~2009a,2009b; Gilmour et al.~2009). These observations suggest that increased nuclear activity may be triggered in such environments and that AGN-driven outflows may play a role in suppressing star formation within galaxies during cluster assembly. Indeed the increased dynamical friction within groups and their low relative velocity dispersions make them conducive to galaxy interactions which can trigger such activity (Hickson 1997; Canalizo \\& Stockton 2001) and recent hydrodynamical simulations suggest that merger-triggered AGN feedback can have a profound effect on the gas content and star formation activity of their host galaxies (Hopkins et al.~2007; Somerville et al.~2008).  Since the mass density of virialized structures increases with redshift, mergers are expected to play an even greater role in the group environment in the past.  Therefore, if galaxy interactions and subsequent AGN feedback are driving a significant portion of the preprocessing found in intermediate density environments, we may expect to find an overdensity of AGN in high redshift groups in the early stages of hierarchical cluster formation.  In this Letter we report on \\emph{Chandra} observations of one such system, the Cl0023+0423 (hereafter Cl0023) four-way group merger at $z=0.84$ (Lubin et al.~2009a).  The Cl0023 structure consists of four interacting galaxy groups which simulations suggest are the direct progenitors of a future massive cluster.  As such, the system provides a unique look at the level of processing and evolution already underway in the group environment prior to cluster assembly.  To search for evidence of triggered nuclear activity within the Cl0023 structure, we present the number counts of X-ray point sources detected in a field covering the entire system, as well as a cross-correlation of these sources with our extensive spectroscopic database.  Surprisingly, we find no evidence for an overdensity of X-ray detected point sources in the direction of the Cl0023 groups.  We discuss the implications of this finding on the role of AGN feedback in regulating galaxy evolution in such structures. We also examine possible explanations for the lack of increased nuclear activity in the system.  Throughout this Letter we assume a $\\Lambda$CDM cosmology with $\\Omega_{m} = 0.3$, $\\Omega_{\\Lambda} = 0.7$, and $H_{0} = 70$ $h_{70}$ km s$^{-1}$ Mpc$^{-1}$.   \\begin{figure*} \\epsscale{1.1} \\plotone{./f1.astroph.eps} \\caption{(\\emph{left}) Adaptively smoothed density map of color-selected red galaxies in the Cl0023 field.  Three density peaks that correspond with four spectroscopically confirmed galaxy groups are marked. Adapted from Lubin et al.~(2009a).  (\\emph{right}) Adaptively smoothed, ACIS-I image of the Cl0023 field in the soft X-ray band (0.5-2 keV). \\label{fig-densmap}} \\end{figure*}  \\vspace{0.5in} ",
        "conclusions": "Using \\emph{Chandra} imaging of the Cl0023 complex, we have searched for evidence of triggered nuclear activity within a dynamically active system of four galaxy groups in the early stages of cluster formation.   Both the redshift distribution and cumulative number counts of X-ray point sources in the Cl0023 field reveal little evidence to suggest that the system contains X-ray luminous AGN in excess to what is observed in the field population. These results are at odds with previous reports of source excesses on the outskirts of dynamically unrelaxed clusters at high redshift.  They also appear to challange the notion that AGN-driven outflows play a significant role in the preprocessing observed in galaxy groups and environments of moderate overdensity relative to the field. If preprocessing is underway in the Cl0023 system, our observations suggest that powerful (quasar mode) nuclear activity is not the predominant mechanism quenching star formation and driving the evolution of Cl0023 galaxies.   Of course we cannot rule out a population of low-luminosity AGN powering ``radio mode'' feedback (Croton et al.~2006) in the Cl0023 complex as our observations are only sensitive to moderate luminosity Seyferts and QSOs.  We are currently analyzing \\emph{Very Large Array} (\\emph{VLA}) 20-cm observations of Cl0023 to search for such a population and expect to present a full radio study of the system in a forthcoming paper (L.~M.~Lubin et al.~2009b, in preparation).  Our current findings are in stark contrast to the overdensity of AGN recently detected in similar \\emph{Chandra} observations of the Cl1604 supercluster at $z=0.9$, where we find a population of Seyferts associated with an unrelaxed cluster and two rich groups (Kocevski et al.~2009a, 2009b).  However, the galaxy populations of these groups differ in significant ways from those of the Cl0023 system.  The Cl1604 groups tend to have higher velocity dispersions and more evolved galaxy populations than the Cl0023 groups, as indicated by their average SFRs and morphological fractions (Gal et al.~2008; Lubin et al.~2009a).  Previous observations of Cl0023 galaxies found them to be predominately late-type systems (75\\%; Lubin et al.~1998) with substantial amounts of ongoing star formation\\footnote{This is consistent with the galaxy properties of high redshift groups with similar velocity dispersions (e.g.~Poggianti et al.~2006)} (Postman, Lubin, Oke 1998; Lubin et al.~2009a), whereas the hosts of the Cl1604 AGN tend be bulge-dominated, post-starburst galaxies which show signs of recent or ongoing galaxy interactions.   Therefore, while Cl0023 contains galaxies which have the gas necessary to fuel nuclear activity, it apparently lacks the bulge-dominated and massive early-type hosts in which powerful AGN have been shown to reside (Kauffmann et al.~2003).  A likely explanation for the absence of luminous AGN in the Cl0023 groups is that the system lacks galaxies with sufficiently massive nuclear black holes required to power such activity.  It has previously been shown that the bulge-dominated S0 population in clusters and groups builds up over time at the expense of the spiral population and that this morphological evolution is more pronounced in lower mass systems (Poggianti et al.~2009).  There is also evidence that these galaxies are typically more massive than their suspected progenitors (Dressler et al.~2009), suggesting they experience growth in their stellar bulges while in overdense environments, possibly via a centrally concentrated burst of star formation (Dressler et al.~1999). Given the correlation between bulge mass and central black hole mass (Gebhardt et al.~2000), we would expect similar growth in galactic nuclei over the same period.  Therefore, if disruptive AGN-driven outflows play a role in quenching star formation in groups, as has been suggested, it may only become an important factor in the preprocessing of galaxy populations during a later stage in the evolution of such groups and structures, when sufficiently massive galaxies (and nuclear black holes) have built up, but prior to hydrodynamical processes within clusters stripping them of their gas reservoirs.  Further observations of a larger sample of systems in the early stages of cluster formation, with a variety of velocity dispersions and morphological fractions, will be required to test this scenario. In the mean time, we are planning additional spectroscopic follow-up of the Cl0023 groups targeting the radio bright population as well as the remaining X-ray point sources that currently lack redshifts. This will give us a greater spectroscopic completeness of X-ray luminous AGN in the Cl0023 field, which will enable us to test our current findings and should allow us to better discern the prevalence of powerful nuclear activity during cluster formation."
    },
    "0910/0910.2234_arXiv.txt": {
        "abstract": "We present a physical model for the origin of \\zsim 2 Dust-Obscured Galaxies (DOGs), a class of high-redshift ULIRGs selected at 24 \\micron \\ which are particularly optically faint ($F_{\\rm 24\\mu m}/F_R>1000$). By combining $N$-body/SPH simulations of high redshift galaxy evolution with 3D polychromatic dust radiative transfer models, we find that luminous DOGs (with \\stf$\\ga 0.3 $ mJy at $z\\sim 2$) are well-modeled as extreme gas-rich mergers in massive ($\\sim 5\\times10^{12}-10^{13}$ \\msunend) halos, with elevated star formation rates ($\\sim 500-1000$ \\msunyrend) and/or significant AGN growth ($\\dot{M}_{\\rm BH} \\ga 0.5$ \\msunyrend), whereas less luminous DOGs are more diverse in nature.  At final coalescence, merger-driven DOGs transition from being starburst dominated to AGN dominated, evolving from a ``bump'' to a power-law shaped mid-IR (IRAC) spectral energy distribution (SED). After the DOG phase, the galaxy settles back to exhibiting a ``bump'' SED with bluer colors and lower star formation rates.  While canonically power-law galaxies are associated with being AGN-dominated, we find that the power-law mid-IR SED can owe both to direct AGN contribution, as well as to a heavily dust obscured stellar bump at times that the galaxy is starburst dominated. Thus power-law galaxies can be either starburst or AGN dominated. Less luminous DOGs can be well-represented either by mergers, or by massive ($M_{\\rm baryon} \\approx 5\\times10^{11}$\\msun) secularly evolving gas-rich disc galaxies (with SFR $\\ga 50$ \\msunyrend).  By utilising similar models as those employed in the SMG formation study of Narayanan et al. (2010), we investigate the connection between DOGs and SMGs.  We find that the most heavily star-forming merger driven DOGs can be selected as Submillimetre Galaxies (SMGs), while both merger-driven and secularly evolving DOGs typically satisfy the \\bzk \\ selection criteria.  The model SEDs from the simulated galaxies match observed data reasonably well, though Mrk 231 and Arp 220 templates provide worse matches. Our models provide testable predictions of the physical masses, dust temperatures, CO line widths and location on the \\magorrian \\ relation of DOGs. Finally, we provide public SED templates derived from these simulations. ",
        "introduction": "\\label{section:introduction}    Redshift \\zsim 2 is a rich epoch for understanding galaxy formation and evolution. During this time period, the bulk of the cosmic stellar mass was assembled \\citep{dic03,rud06}, and the star formation rate and quasar space densities are both near their peaks \\citep{bou04,hop04,hop07,sha96}.  In the local Universe, infrared (IR) selection of galaxies is an efficient means of selecting galaxies undergoing significant star formation (e.g., Arp 220), and possibly hosting optically-obscured active galactic nuclei (AGN; e.g., Mrk 231, NGC 6240). By analogy, at higher redshifts, IR selection techniques select galaxies which are major contributors to the cosmic star formation rate density, far-IR background, and progenitors of the present-day massive galaxy population. Indeed, ultraluminous infrared galaxies (ULIRGs; \\lir $>$ 10$^{12}$ \\lsunend) at \\zsim 2 appear to contribute substantially to the infrared luminosity density \\citep[]{per05,cap07,red08,hop10,hop10b}, and the \\z$>$0.7 star formation rate density \\citep{lef05,shi09}.   A well-studied sample of galaxies selected for their FIR properties at \\zsim 2 are Submillimetre Galaxies (SMGs), chosen in blank field surveys for their redshifted cool dust emission at 850 \\micron \\ above \\sef $\\ga$ 5 mJy. These galaxies are in a transient starburst phase \\citep[e.g.,] []{swi04,men07,you08b,mic09}, host rapidly growing supermassive black holes \\citep[e.g.,] []{ale05a,ale08}, and reside in extremely massive $\\sim$5$\\times$10$^{12}$ \\msun halos \\citep{bla04}. While these galaxies represent an important sub-population of infrared-luminous galaxies at the epoch of peak galaxy formation, there is some concern that a selection at 850 \\micron \\ may be missing significant populations of high-redshift ULIRGs with warmer dust temperatures \\citep[e.g.,] []{bus09,cha09,cas09,hua09,you09b}. As such, alternative selection techniques for identifying \\zsim 2 ULIRGs have become desirable.  With the launch of the {\\it Spitzer Space Telescope}, surveys at 24 \\micron \\ have uncovered a population of ULIRGs at \\zsim 2 \\cite[e.g.,] [ and references therein]{rig04,don07,yan07,far08,soi08,lon09,hua09}.  Recently, \\citet[][ hereafter, D08; see also \\citet{fio08}]{dey08} presented an efficient means of identifying high-redshift ULIRGs which are both mid-infrared luminous and optically faint. Specifically, by imposing a nominal selection criteria of $R$-[24] $>$ 14 (roughly corresponding to a flux density ratio \\stf/\\sr $>$ 1000; formally \\stf/\\sr $> 960$) and \\stf $>$ 0.3 mJy, D08 found a sample of ULIRGs with a relatively narrow redshift distribution (centered at $<$\\z$>$=1.99 with $\\sigma_z$=0.5) which exhibit heavy reddening of the intrinsic UV flux and strong rest-frame 8 \\micron \\ emission \\citep{hou05,fio08}. D08 refer to these galaxies as Dust-Obscured Galaxies (DOGs)\\footnote{We note that while this paper nominally focuses on 24 \\micron-selected galaxies which additionally have \\tfr $> 1000$, the results of this paper are generally applicable to the 24 \\micron-bright (\\stf $> 0.3 $mJy) \\zsim 2 ULIRG population.}.    DOGs exhibit a range of luminosities, though AGN-dominated galaxies tend to preferentially dominate the bright end of the 24 \\micron \\ flux density distribution \\citep{bra06,wee06b,wee06a,dey08,fio08,fio09,sac09}. The DOGs with the highest 24 \\micron \\ flux densities (\\stf $>$ 0.8 mJy) have bolometric infrared luminosities (\\lir) $\\sim$ 10$^{13}$ \\lsun \\citep{bus09b,tyl09}.  As a population, the subset of \\stf$>0.3$ mJy galaxies which are also DOGs contribute around a quarter of the total infrared luminosity density at \\zsim 2 \\citep{dey08}, and constitute $\\sim$60\\% of the total contribution by \\zsim 2 ULIRGs. It is clear that DOGs are a cosmologically significant population of galaxies with diverse observational characteristics.  Broadly, 24 \\micron-selected ULIRGs at \\zsim 2 (including DOGs) appear to fall into two categories based on their mid-IR IRAC SEDs: those with a power-law shape and those which exhibit a bump at observed $\\lambda \\sim 5$ \\micron \\ corresponding to the 1.6 \\micron \\ rest frame stellar photospheric bump (hereafter, these are referred to as PL galaxies and bump galaxies, respectively). Bump galaxies are thought to be dominated by star formation, with the bump (at rest-frame 1.6 \\micron) originating from starlight.  PL galaxies are thought to have their mid-IR SED dominated by AGN continuum emission \\citep{wee06a,don07,yan07,mur09}.  In support of this scenario, PL galaxies have SEDs which are reasonably represented by that of a Mrk 231 template \\citep{bus09b,tyl09}, and tend to show relatively compact optical morphologies \\citep[$R_e\\sim 1-6$ kpc; ][]{mel08,mel09,bus09}. On the other hand, bump galaxies are typically polycyclic aromatic hydrocarbon (PAH)-rich, show PAH equivalent widths consistent with being star formation dominated \\citep{yan05,saj07,far08,des09,hua09}, and are more often morphologically resolved into multiple components \\citep{bus09}. However, \\citet{fio08,fio09} have argued that even these less luminous galaxies may harbor heavily obscured AGN.  DOGs are massive galaxies, with stellar masses $M_\\star >$ 10$^{10-11}$ \\msun \\citep{bus09,lon09,hua09}, and cluster in group-sized halos of order \\citep[$M_{\\rm DM} \\approx 10^{12-13}$ \\msunend;][]{bro08}. The clustering is luminosity-dependent, with more luminous DOGs inhabiting more massive halos. Since the more luminous DOGs also tend to exhibit power-law SEDs, it follows that PL DOGs cluster more strongly than the lower luminosity bump DOGs.   Although much progress has been made in understanding DOGs in a relatively short time period, a myriad of fundamental questions regarding their physical nature exist. Are DOGs preferentially isolated galaxies or mergers (or both)? Does the \\tfr \\ criterion select galaxies at a particular evolutionary point? How are bump DOGs and PL DOGs related - are they distinct galaxy populations, or possibly related via an evolutionary sequence? How are DOGs related to other coeval high-redshift populations (e.g., the SMGs, \\bzk \\ galaxies, etc.)?  Numerical simulations can offer complementary information to the observations in hand, and provide insight into the aforementioned questions. In particular, by coupling radiative transfer modeling with hydrodynamic simulations of galaxy evolution, one can precisely mimic local and high-redshift observational selection functions, and directly relate observed trends to physical conditions in the model galaxies \\citep[e.g.,] []{jon06a,jon06b,jon10,lot08,nar10a,you09}.  Previously, we have utilised similar methods to investigate the formation and evolution of high-redshift Submillimetre Galaxies \\citet[][C. Hayward et al. in prep.]{nar09b,nar10a}. Here, we build upon these efforts by utilising similar models as those employed in \\citet{nar09b} and \\citet{nar10a}.  This will allow us not only to investigate a potential formation mechanism for DOGs, but to explore the potential connection between \\zsim 2 ULIRGs selected for their 24 \\micron \\ brightness, and those selected in the submillimetre.  The paper is organised as follows. In \\S~\\ref{section:methods}, we detail our numerical methods. In \\S~\\ref{section:dogformation}, we describe the formation of \\zsim 2 DOGs, and explain the evolution of the 24/R ratio and the relationship between bump and PL DOGS. In \\S~\\ref{section:physicalnature}, we discuss the physical requirements necessary to form a DOG, and whether mergers are necessary; In \\S~\\ref{section:othergalaxies}, we detail the relationship between DOGs, SMGs and \\bzk \\ galaxies. In \\S~\\ref{section:observations}, we relate our model to existing observations, and make testable predictions. In \\S~\\ref{section:discussion}, we provide discussion. We conclude in \\S~\\ref{section:conclusions}.  Throughout this paper, we utilise a fiducial DOG selection criteria \\tfr \\ $>$ 1000 and \\stf $>$ 300 \\microjy, consistent with the DOGs sample in D08. We note again that while we focus on these 24 \\micron \\ selected galaxies which are optically faint (DOGs), the results are generally applicable to most 24 \\micron -selected \\zsim 2 ULIRGs. Throughout, we assume a cosmology with $\\Omega_\\Lambda=0.7$, $\\Omega_m=0.3$ and $h=0.7$.  Finally, we note that the models presented here are not cosmological in nature. As such, all non-evolutionary plots should be taken to represent ranges of expected values and colors, rather than true expected distributions.      \\begin{table*} \\label{table:ICs} \\centering \\begin{minipage}{100mm} \\caption{DOGs are ordered with decreasing halo mass (and grouped by mergers and isolated galaxies, such that the isolated galaxies have the prefix 'i' in their name). Column 1 is the name of the model used in this work.  We emphasise that these models are nearly identical to those employed in the SMG formation study by \\citet{nar10a} and \\citet{nar09b}.  This will facilitate the comparison of these models with SMGs (\\S~\\ref{section:smg}). Columns 2 and 3 give the virial velocity and halo mass of the galaxies. Column 4 is the total baryonic mass of the system. Column 5 is the merger mass ratio. Columns 6 \\& 7 are initial orientations for disc 1, Columns 8 \\& 9 are for disc 2. The orientation angles are with respect to the merger plane (such that $\\theta , \\phi = 0$ would be a coplanar merger). Column 10 lists the initial gas fraction of the simulation.} \\begin{tabular}{@{}cccccccccc@{}} \\hline Model & V$_{\\rm c}$ & $M_{\\rm DM}$ & $M_{\\rm bar}$& Mass Ratio & $\\theta_1$&$\\phi_1$ & $\\theta_2$ & $\\phi_2$ &$f_g$\\\\ &(\\kmsend)&\\msun&\\msun&&&&\\\\ \\hline DOG1 & 500:500 & 1.25$\\times$10$^{13}$:1.25$\\times$10$^{13}$ & 1.1$\\times$10$^{12}$& 1:1 & 30 & 60 & -30 & 45 &0.8,0.6,0.4\\\\ DOG2 & 500:500 &  1.25$\\times$10$^{13}$:1.25$\\times$10$^{13}$ & 1.1$\\times$10$^{12}$& 1:1& 360 & 60 & 150 & 0& 0.8\\\\ DOG3 & 500:500 &  1.25$\\times$10$^{13}$:1.25$\\times$10$^{13}$& 1.1$\\times$10$^{12}$& 1:1 &-109 & -30 & 71 &-30& 0.8\\\\ DOG4 & 500:320 & 1.25$\\times$10$^{13}$:3.4$\\times$10$^{12}$ & 6.9$\\times$10$^{11}$&1:3 & 30 & 60 & -30 & 45& 0.8,0.6,0.4\\\\ DOG5 & 500:320 &  1.25$\\times$10$^{13}$:3.4$\\times$10$^{12}$& 6.9$\\times$10$^{11}$&1:3 & 360 & 60 & 150 & 0 & 0.8\\\\ DOG6 & 500:320 &  1.25$\\times$10$^{13}$:3.4$\\times$10$^{12}$& 6.9$\\times$10$^{11}$&1:3& -109 & -30 & 71 & -30 & 0.8\\\\ DOG7 & 500:225 & 1.25$\\times$10$^{13}$:1.2$\\times$10$^{12}$ & 6.0$\\times$10$^{11}$&1:10 & 30 & 60 & -30 &45 & 0.8,0.6,0.4\\\\ DOG8 & 500:225 & 1.25$\\times$10$^{13}$:1.2$\\times$10$^{12}$ & 6.0$\\times$10$^{11}$&1:10 & 360 & 60 & 150 & 0 & 0.8\\\\ DOG9 & 500:225 & 1.25$\\times$10$^{13}$:1.2$\\times$10$^{12}$ & 6.0$\\times$10$^{11}$&1:10 & -109 & -30 & 71 &-30 & 0.8\\\\ DOG10 & 320:320 & 3.4$\\times$10$^{12}$:3.4$\\times$10$^{12}$ &2.9$\\times$10$^{11}$ & 1:1 & 30 & 60 & -30 &45 &0.8,0.6,0.4\\\\ DOG11 & 320:320 & 3.4$\\times$10$^{12}$:3.4$\\times$10$^{12}$ &2.9$\\times$10$^{11}$ & 1:1 & 360 & 60 & 150 &0 &0.8\\\\ DOG12 & 320:320 & 3.4$\\times$10$^{12}$:3.4$\\times$10$^{12}$ &2.9$\\times$10$^{11}$ & 1:1 &-109 & -30 & 71 &-30& 0.8\\\\ DOG13 & 225:225 & 1.2$\\times$10$^{12}$:1.2$\\times$10$^{12}$ & 1.0$\\times$10$^{11}$& 1:1 & 30 & 60 & -30 &45 &0.8,0.6,0.4\\\\ iDOG1 & 500 & 1.25$\\times$10$^{13}$ &5.5$\\times$10$^{11}$& --- & N/A & N/A&N/A & N/A& 0.8,0.6,0.4\\\\ iDOG2 & 320 & 3.4$\\times$10$^{12}$ & 1.5$\\times$10$^{11}$& --- & N/A & N/A&N/A & N/A &0.8,0.6,0.4\\\\ iDOG3 & 225 & 1.2$\\times$10$^{12}$ &5$\\times$10$^{10}$& ---& N/A &N/A& N/A & N/A &0.8,0.6,0.4\\\\    \\hline \\end{tabular} \\end{minipage} \\end{table*} ",
        "conclusions": "\\label{section:conclusions}  Utilising a combination of polychromatic radiative transfer calculations and hydrodynamic simulations of high-redshift galaxy evolution, we have formulated a physical model for the formation and evolution of \\zsim 2 Dust-Obscured Galaxies. We have utilised the very similar models as those employed in previous studies aimed at investigating the formation and evolution of a similar class of \\zsim 2 ULIRGS: the Submillimetre Galaxy population.  While the uncertainty in the models is dominated by small-scale physics which is unresolved by the simulations, our methodology allows us to make the following general conclusions regarding the origin of DOGs, and their relationship to other high-\\z \\ ULIRGs:  \\begin{itemize}  \\item DOGs are a diverse class of galaxies, ranging from secularly evolving star-forming disc galaxies (forming stars at $\\sim$50-100 \\msunyrend) to extreme gas-rich galaxy mergers forming stars at $\\ga$ 1000 \\msunyrend.  \\item The most luminous DOGs (\\stf $\\ga$ 300 \\microjy) are well represented by mergers of massive ($M_{\\rm DM} \\approx 5 \\times 10^{12}-10^{13}$ \\msunend) galaxies. These galaxies are actively transitioning from being starburst dominated to being AGN dominated. At decreasing 24 \\micron \\ flux densities (\\stf $\\la$ 100 \\microjy), DOGs may either be galaxy mergers, or secularly evolving disc galaxies with more modest ($\\sim$50-100 \\msunyrend) SFRs.  \\item Luminous, merger-driven DOGs naturally transition from having a bump-like mid-IR SED to a power-law (PL) shape, as the dominant power source transitions from star formation to the central AGN. That said, there is an overlap period where star-formation dominated sources can appear as PL galaxies.  \\item The most luminous DOGs assemble both significant stellar masses ($M_\\star \\approx 10^{11}$ \\msunend), as well as contribute toward the growth of supermassive ($M_{\\rm BH} \\approx 10^9$ \\msunend) black holes.  \\item Merger-driven DOGs overlap with the \\zsim 2 Submillimetre Galaxy population. SMGs generally represent the earlier, starburst-dominated phase of the DOGs evolution. This can be tested via the location of SMGs and DOGs on the \\magorrian \\ relation.    \\end{itemize}  In advance of {\\it Herschel}, {\\it JWST} and ALMA, we provide the following testable predictions for our model of DOG formation:  \\begin{itemize}  \\item We quantify the expected range of black hole, \\htwo, stellar, and halo masses for luminous (\\stf $>$ 300 \\microjy) DOGs as well as SMGs.  \\item We provide dust temperatures for DOGs selected at various 24 \\micron \\ flux density limits.  \\item We detail the location of SMGs and bright (\\stf $>$ 1 mJy) DOGs on the \\magorrian \\ relation as a test for the modeled evolutionary scenario that SMGs evolve into PL DOGs.  \\item We provide a prediction for the CO line widths from DOGs, and suggest that they will be of order the largest observed line widths at high redshift, comparable to \\zsim 2 SMGs.  \\end{itemize}  Finally, we provide model SED templates for \\z=2 DOGs as a function of total infrared luminosity (\\lir = 1-1000 \\micron). These are provided publicly, along with an IDL script for easily reading in the templates at http://www.cfa.harvard.edu/~dnarayan/DOG\\_SED\\_Templates.html"
    },
    "0910/0910.4718_arXiv.txt": {
        "abstract": "The Array for Microwave Background Anisotropy (AMiBA) is a radio interferometer for research in cosmology, currently operating 7 0.6m diameter antennas co-mounted on a 6m diameter platform driven by a hexapod mount. AMiBA is currently the largest hexapod telescope. We briefly summarize the hexapod operation with the current pointing error model. We then focus on the upcoming 13-element expansion with its potential difficulties and solutions. Photogrammetry measurements of the platform reveal deformations at a level which can affect the optical pointing and the receiver radio phase. In order to prepare for the 13-element upgrade, two optical telescopes are installed on the platform to correlate optical pointing tests. Being mounted on different locations, the residuals of the two sets of pointing errors show a characteristic phase and amplitude difference as a function of the platform deformation pattern. These results depend on the telescope's azimuth, elevation and polarization position. An analytical model for the deformation is derived in order to separate the local deformation induced error from the real hexapod pointing error. Similarly, we demonstrate that the deformation induced radio phase error can be reliably modeled and calibrated, which allows us to recover the ideal synthesized beam in amplitude and shape of up to 90\\% or more. The resulting array efficiency and its limits are discussed based on the derived errors. ",
        "introduction": "The Array for Microwave Background Anisotropy (AMiBA) is a forefront radio interferometer for research in cosmology. The project is led by the Academia Sinica, Institute of Astronomy and Astrophysics (ASIAA), Taiwan, with major collaborations with National Taiwan University, Physics Department (NTUP), Electrical Engineering Department (NTUEE), and the Australian Telescope National Facility (ATNF). Contributions also came from the Carnegie Mellon University and NRAO. As a dual-channel 86-102 GHz interferometer array of up to 19 elements, AMiBA is designed to have full polarization capabilities, sampling structures greater than $2^{\\prime}$ in size. The AMiBA targets the distribution of clusters of galaxies via the Sunyaev-Zel'dovich (SZ) Effect$^1$. It will also measure the Cosmic Microwave Background (CMB) temperature anisotropies$^2$ on scales which are sensitive to structure formation scenarios of the Universe. AMiBA is sited on Mauna Loa, Big Island, Hawaii, at an elevation of 3,425m, currently operating the initial phase with 7 antennas of 0.6m diameter in a compact configuration, Fig.\\ref{front_hexpol0_bw}. The correlator and the receivers with the antennas are co-mounted on a fully steerable 6m diameter platform controlled by a hexapod mount. The AMiBA hexapod mount with its local control system was designed and manufactured by Vertex Antennentechnik GmbH, Duisburg, Germany. The 0.6m diameter Cassegrain antenna$^3$ field-of-view (fov) is $\\sim 22^{\\prime}$ at Full Width Half Maximum (FWHM). The synthesized beam resolution in the hexagonally compact configuration (0.6m, 1.04m and 1.2m baselines) is $\\sim 6^{\\prime}$. This sets our target pointing accuracy to about $\\sim 0.6^{\\prime}$. \\\\ In parallel to the current 7-element observation program, AMiBA is undergoing a 13-element expansion with 1.2m diameter Cassegrain antennas (fov $\\sim 11^{\\prime}$). The resulting synthesized beam of about $\\sim 2^{\\prime}$ requires then ideally an improved pointing accuracy of about $\\sim 0.2^{\\prime}$. The platform photogrammetry tests$^{4,5}$ have revealed local deformations which might affect individual antenna pointings (including the optical telescope pointing) and their radio phases for the longer baselines. It therefore seems desirable to include such a correction in the current pointing error model and in the phase correction scheme. \\\\ The goal of this paper is to investigate a refined pointing error and phase correction model for the AMiBA expansion phase. The adopted  approach can be of general interest since a platform deformation for a hexapod of this size is likely to be a generic problem depending on the level of accuracy which is needed. The paper is organized as following: Section \\ref{7element} summarizes the pointing error model used for the currently operating 7-element compact configuration. Section \\ref{13element} presents the approach taken to improve the pointing accuracy and the phase correction for the 13-element expansion and it also discusses its limitations. Previous AMiBA status reports and the current observations are described in a series of papers.$^{6,7,8,9}$   \\begin{figure} \\begin{center} \\includegraphics[scale=0.95]{front_hexpol0_archive.eps} \\caption{\\label{front_hexpol0_bw} Front view of the AMiBA. Installed are 7 antennas in compact configuration, giving 0.6m, 1.04m and 1.2m baselines. Free receiver holes in the platform allow for different array configurations and for the expansion phase with 13 antennas. The 1st 8 inch refractor (OT1) for optical pointing is attached to the black bracket below the lowermost antenna at a distance of 1.4m from the platform center. A second optical telescope (OT2) is located at 90$^{\\circ}$ to the left hand side at the same radius.} \\end{center} \\end{figure}   \\begin{figure} \\begin{center} \\includegraphics[scale=0.8]{ot1_ot2_bw.eps} \\caption{\\label{ot1_ot2__bw}OT1 (in the back) and OT2 (in the front, lower left hand side corner), mounted with a $90^{\\circ}$ phase shift in between them at a radius of 1.4m. Clearly seen are also the reflecting photogrammetry targets.} \\end{center} \\end{figure} ",
        "conclusions": "We have demonstrated that the local platform deformation error can be isolated from the pointing error with the help of two OTs. The extracted deformation error is in the range $-1'$ to $+1'$ which is consistent with the photogrammetry data. The verification of an improved pointing accuracy is, however, non-trivial, because the platform deformation component will always be present on the ccd images. Therefore, one always has to rely on the modeling. Aiming for a pointing accuracy of $\\sim 0.5'$ rms ($\\eta_{abs}\\sim 0.76$) seems reasonable. Even smaller pointing errors improve the efficiency only marginally. Similarly, the platform deformation induced phase error is corrected with the photogrammetry data, and the synthesized beam amplitude is restored to 90\\% or more. \\\\ Finally, one has to ask whether the separation of platform and pointing error is needed at all. Assuming that the errors are not separated and the derived pointing errors are then reduced to zero, one still ends up with an effective pointing error ($p_{az}$ and $p_{el}$ in Eq.(\\ref{delta_abs})) in the range $[-1,1]$ due to the over-/under-compensated platform deformation error. $\\delta_{abs}$ can therefore not be further reduced and the efficiency is limited to 67\\%. For large pointing errors ($>1'$) the separation becomes unnecessary. Going beyond $\\eta_{abs}\\sim 0.67$ requires a better effective pointing and a separation of the error components. In any case, a  phase correction model is needed to restore the synthesized beam."
    },
    "0910/0910.4835_arXiv.txt": {
        "abstract": "{The e-VLBI technique offers a unique opportunity for users to probe the milliarcsecond (mas) scale structure of unidentified radio sources, and organise quick follow-up observations in case of detection. Here we report on e-EVN results for a peculiar radio source that has been suggested to act as a gravitational lens. However the lensing galaxy has not been identified in the optical or the IR bands so far. Our goal was to look for an active galactic nucleus (AGN) in this suspected dark lens system. The results indicate strong AGN activity, and rule out the possibility that the radio source itself is gravitationally lensed.\\ }  \\FullConference{Science and Technology of Long Baseline Real-Time Interferometry: \\\\ The 8th International e-VLBI Workshop, EXPReS09\\\\ June 22 - 26 2009\\\\ Madrid, Spain}  \\begin{document} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.3412_arXiv.txt": {
        "abstract": "{ We present synthetic broad-band photometric colors of a late-type giant located close to the RGB tip ($T_{\\rm eff}\\approx3640$\\,K, $\\log g=1.0$ and ${\\rm [M/H]}=0.0$). Johnson-Cousins-Glass \\emph{BVRIJHK} colors were obtained from the spectral energy distributions calculated using 3D hydrodynamical and 1D classical stellar atmosphere models. The differences between photometric magnitudes and colors predicted by the two types of models are significant, especially at optical wavelengths where they may reach, e.g., $\\Delta V\\approx0.16$, $\\Delta R\\approx0.13$ and $\\Delta (V-I)\\approx0.14$, $\\Delta (V-K)\\approx0.20$. Differences in the near-infrared are smaller but still non-negligible (e.g., $\\Delta K\\approx 0.04$). Such discrepancies may lead to noticeably different photometric parameters when these are inferred from photometry (e.g., effective temperature will change by $\\Delta T_{\\rm eff}\\approx60$\\,K due to difference of $\\Delta (V-K)\\approx0.20$). ",
        "introduction": "Late-type giants are important tracers of intermediate age and old stellar populations in the Galaxy and beyond. Thus, the availability of reliable stellar atmosphere models is of utmost importance for understanding the structure and evolution of these stars and their host populations. The advent of the 3D hydrodynamical codes allowed to accommodate a more realistic treatment of non-stationary phenomena (e.g., convection) in the stellar atmosphere modeling and thus it is natural to expect that new 3D hydrodynamical models may provide important insights about observable properties  and the interior structures of late-type giants. Indeed, recent work of \\citet{CAT07}, \\citet{C08} demonstrates that there are significant differences in the spectral line strengths predicted by the 3D hydrodynamical and 1D classical stellar atmosphere models. This leads to substantial discrepancies between the elemental abundances of various chemical species derived with the classical 1D and 3D hydrodynamical stellar atmosphere models \\citep[see][for details]{CAT07}.  Still, the extent of differences in the global spectral properties predicted by the 1D classical and 3D hydrodynamical models is largely unknown. In one of our previous studies we have used a simplified approach to show that significant differences can be expected in the photometric colors predicted by the two types of models \\citep{KHL05}. Whether these conclusions would also hold with the photometric colors based on the results of full 3D spectral synthesis still had to be confirmed.  In this study we extend the previous work and calculate synthetic photometric colors of a late-type giant from the full 3D spectral energy distributions. For this purpose we use a model of a late-type giant which is significantly cooler than those considered by Collet et al. (2007), with its atmospheric properties typical to those of solar-metallicity late-type giants located close to the RGB tip.   \\begin{figure*}[t!] \\begin{centering} \\resizebox{10cm}{!}{\\includegraphics[clip=true]{kucinskas_fig02.eps}} \\caption{\\footnotesize Differences between the photometric magnitudes of a late-type giant predicted by 3D hydrodynamical and 1D classical models: ${\\rm 3D}-{\\rm 1D}$ (black solid line) and ${\\rm 3D}-\\mD$ (red/gray dashed line).} \\label{colors-BVRIJHK} \\end{centering} \\end{figure*} ",
        "conclusions": "The results obtained allow us to conclude that convection indeed plays an important role in defining the intrinsic atmospheric structures and observed properties of late-type giants. Specifically, we find that spectral energy distributions and photometric colors of a late-type giant produced with 3D hydrodynamical and 1D classical stellar atmosphere models are substantially different. Differences in photometric magnitudes and colors are considerably larger than typical photometric errors (e.g., $\\Delta V\\approx0.16$, $\\Delta (V-K)\\approx0.20$). These differences may result in further discrepancies, for instance, in the photospheric parameters derived from photometric colors (e.g., a difference of $\\Delta (V-K)\\approx0.20$ will change the estimated effective temperature by $\\Delta T_{\\rm eff}\\approx60$\\,K). Obviously, this may have direct consequences to any photometric work that relates to late-type giants and thus once again stresses the importance of 3D hydrodynamical model atmospheres to be used in the interpretation of observational data."
    },
    "0910/0910.4697_arXiv.txt": {
        "abstract": "We develop an algorithm of separating the $E$ and $B$ modes of the CMB polarization from the noisy and discretized maps of Stokes parameter $Q$ and $U$ in a finite area. A key step of the algorithm is to take a wavelet-Galerkin discretization of the differential relation between the $E$, $B$ and $Q$, $U$ fields. This discretization allows derivative operator to be represented by a matrix, which is exactly diagonal in scale space, and narrowly banded in spatial space. We show that the effect of boundary can be eliminated by dropping a few DWT modes located on or nearby the boundary. This method reveals that the derivative operators will cause large errors in the $E$ and $B$ power spectra on small scales if the $Q$ and $U$ maps contain Gaussian noise. It also reveals that if the $Q$ and $U$ maps are random, these fields lead to the mixing of the $E$ and $B$ modes. Consequently, the $B$ mode will be contaminated if the powers of $E$ modes are much larger than that of $B$ modes. Nevertheless, numerical tests show that the power spectra of both $E$ and $B$ on scales larger than the finest scale by a factor of 4 and higher can reasonably be recovered, even when the power ratio of $E$- to $B$-modes is as large as about 10$^2$, and the signal-to-noise ratio is equal to 10 and higher. This is because the Galerkin discretization is free of false correlations, and keeps the contamination under control. As wavelet variables contain information of both spatial and scale spaces, the developed method is also effective to recover the spatial structures of the $E$ and $B$ mode fields. ",
        "introduction": "The scalar component of primordial perturbations of the universe can be detected by the maps of temperature fluctuations of Cosmic Microwave Background Radiation (CMBR), while the tensor component of the primordial perturbations has to be probed by the maps of the Stokes parameter $Q$ and $U$ of the linear polarization of the CMBR. A tensor field generally contains electric-like $E$-mode and magnetic-like $B$-modes. In the linear regime, vortical mode of primordial perturbations do not grow during the clustering of density field, and therefore, the perturbed field initially has to be curl-free. That is, the primordial perturbations can only yield the $E$-mode, but not $B$-mode of the CMBR polarization field. On the other hand, $B$-mode perturbations can be produced by gravitational waves. Therefore, extracting the $B$-mode information from CMBR polarization maps is crucial to verify the existence of gravitational wave background produced at the inflationary epoch. Moreover, gravitational lensing of clusters and hot electron scattering of reionization would be able to yield both $E$- and $B$-modes. To study these problems a sharp decomposition of $E$- and $B$-modes from $Q$ and $U$ maps is required.  If both $Q$ and $U$ maps are available over the whole sky, one can find the whole sky maps of $E$- and $B$-modes with the spherical harmonic decomposition, because the relation between the maps of ($Q$, $U$) and ($E$, $B$) in the space spanned by bases of spin two harmonics is local (Kamionkowsky et al. 1997; Zaldarriaga \\& Seljak 1997). However, the observed maps cannot be global; it is always limited by the contamination of our galaxy and other foreground sources. The relation between the maps of ($Q$, $U$) and ($E$, $B$) in physical space contains the Laplace operator, and therefore, it is non-local. The $E/B$ decomposition with the spatially-limited maps of $Q$ and $U$ will not be unique if we lack of information of the polarization and its derivative on the boundary of the maps.  For noiseless samples, the problem of uniqueness would be solved by constructing orthogonal modes with window functions to fit the requirements of boundary conditions (Lewis et al. 2002; Bunn et al. 2003; Smith 2006; Smith \\& Zaldarriaga 2007). It is, however, similar to the domain (or windowed) Fourier analysis. The result will not be useful to study the structures in physical space (e.g. Chiueh \\& Ma 2002).  The other challenge caused by the derivative operator is because the $Q$ and $U$ maps are discrete. Mathematically, the derivative operators $\\partial_x$ or $\\partial^2_x$ are continuous linear operators mapping functions defined in Hilbert space, while the observed samples $Q$ and $U$ actually are defined in space spanned by base $v_i$, $i\\in I$, which is a set of finite indices. This difference leads to large numerical errors when the discrete maps are noisy.  The last, but not least, problem is from the smallness of $B$-modes. On the scale of one degree order, the power of $B$-mode caused by gravitational waves at inflationary epoch is less than that of $E$-modes by a factor of at least 10$^2$. As the maps of $Q$ and $U$ are random fields, the variance of the random fields will lead to the mixing of $E$- and $B$-modes. Consequently, $E$- and $B$-modes would be contaminated from each other. Therefore, it is difficult to recover the power of $B$-mode if the powers of $E$ modes are much larger than that of $B$ modes.  In this paper, we develop an algorithm of the $E/B$ decomposition based on the discrete wavelet transform (DWT) analysis, which is a compromise between the decompositions in physical space and scale space. The DWT analysis of the CMBR temperature fluctuation maps has attracted much attention in the last decade (Pando et al. 1998; Sanz et al. 1999; Mukherjee et al. 2000). Besides these points, we especially take the advantage of the so-called wavelet-Galerkin discretization (e.g. Louis et al. 1997), which is to approximate derivative operator to a matrix in space spanned by wavelet bases. For some available wavelets, the matrixes are exactly diagonal in scale-space, and narrowly banded in spatial space. This made the uncertainties from boundary, noises and variances are under control. We will study the conditions, under which the information of small $B$-mode can approximately be extracted from noisy maps of $Q$ and $U$.  The paper is organized as follows. Section 2 presents the method of the $E/B$ separation in the DWT space. Section 3 tests the DWT algorithm with samples with known spatial structures. We show that the method effectively to suppresses the uncertainties from boundary effect and noise. It is also effective to identify spatial structures of $E$ and $B$ fields. Section 4 presents the tests for samples of Gaussian random field. The effect of the variance of Gaussian random field is analyzed, especially the problem of the mixing of $E$- and $B$-modes. Section 5 addresses the effectiveness of the wavelet-Galerkin discretization. Finally conclusions are given in Section 6. The DWT representation of derivative operators are given in the Appendix. ",
        "conclusions": "The algorithm developed in this paper can be summarized as follows \\begin{itemize} \\item From observed noisy and discrete maps $Q(x,y)$ and $U(x,y)$ we calculate their DWT maps $Q_{l_1,l_2}$ and $U_{l_1,l_2}$ on the finest scale ${\\bf j}=(J,J)$, which is given by the resolution. \\item Using eqs.(16) and (17) we decompose $Q_{l_1,l_2}$ and $U_{l_1,l_2}$ into $\\mathbb{E}_{l_1,l_2}$ and $\\mathbb{B}_{l_1,l_2}$. \\item Using eqs.(19) and (20) we calculate WFCs $\\tilde{\\epsilon}^E_{\\bf j,l}$ and $\\tilde{\\epsilon}^B_{\\bf j,l}$ on scales $(j_1,j_2)$ and $j_1,j_2 \\leq J$. \\item Using the WFC maps $\\tilde{\\epsilon}^E_{\\bf j,l}$  and $\\tilde{\\epsilon}^B_{\\bf j,l}$ we calculate the DWT power spectra by eqs.(30) and (31). \\item We identify spatial structures with maps of $\\mathbb{E}_{l_1,l_2}$ and $\\mathbb{B}_{l_1,l_2}$. \\end{itemize}  With this algorithm, it is possible to recover the power spectrum of $B$-mode random fields from noisy Stokes parameter maps $Q$ and $U$ when the power ratio $E/B$ is as high as $10^2$, and the S/N is equal to or higher than 10. For samples with given structures, the $B$-mode structure can also be identified when the power ratio $E/B$ is equal to 10$^2$. Besides power spectrum, the DWT variables of SFCs ($\\epsilon^E_{\\bf j,l}$, $\\epsilon^B_{\\bf j,l}$) and WFCs ($\\tilde{\\epsilon}^E_{\\bf j,l}$, $\\tilde{\\epsilon}^B_{\\bf j,l}$) can be used for high order statistics, such as high order moments, scale-scale correlation, cross correlation between the $E$ and $B$ and other maps.  With the DWTs, one can construct orthogonal, divergence-free vector wavelets. It has been used for a local analysis of the velocity field of incompressible turbulence (Urban 1995; Kishida et al. 1999; Albukrek et al. 2002). The divergence-free $B$ field is similar to a 2-D velocity field of turbulence of incompressible fluid (e.g. Pina 1998). Therefore, it would be valuable to further study the DWT $E/B$ decomposition with the divergence-free vector wavelets."
    },
    "0910/0910.1598_arXiv.txt": {
        "abstract": "Understanding the details of how the red sequence is built is a key question in galaxy evolution. What are the relative roles of gas-rich vs. dry mergers, major vs. minor mergers or galaxy mergers vs. gas accretion? In a recent paper \\citep{Wild:2009p2609}, we compare hydrodynamic simulations with observations to show how gas-rich major mergers result in galaxies with strong post-starburst spectral features, a population of galaxies easily identified in the real Universe using optical spectra. Using spectra from the VVDS deep survey with $<z>=0.7$, and a principal component analysis technique to provide indices with high enough SNR, we find that 40\\% of the mass flux onto the red-sequence could enter through a strong post-starburst phase, and thus through gas-rich major mergers. The deeper samples provided by next generation galaxy redshift surveys will allow us to observe the primary physical processes responsible for the shut-down in starformation and build-up of the red sequence. ",
        "introduction": "Recent observations have revealed that since a redshift of around unity the total mass of stars living in red sequence galaxies has increased by a factor of two \\citep[e.g.][]{2004ApJ...608..752B}. At the same time, the stellar mass density of the blue sequence has remained almost constant. The interpretation is that some blue galaxies migrate onto the red sequence after the quenching of their star formation, whilst the remainder continue to form new stars \\citep[e.g.][]{2007ApJ...665..265F}.  \\citet{2007A&A...476..137A} measured the net mass flux which has taken place from the blue sequence to the red sequence. This amounts to $9.8\\times10^{-3}$M$_\\odot$/yr/Mpc$^3$, or about $1.4\\times10^4$\\,M$_\\odot$/yr in the VIMOS VLT DEEP Survey (VVDS) volume. Recent star formation history can be used as a tool to identify galaxies as they enter the red sequence. For typical observable times of post-starburst features in VVDS optical galaxy spectra of $\\sim0.35-0.6$\\,Gyr, this could comprise, for example, a few tens of galaxies of stellar mass \\logm$=10.5$ in the VVDS survey. In these proceedings I will summarise the work presented in detail in \\citet{Wild:2009p2609} and ask whether observations of post-starburst galaxies are consistent with the growth of the red-sequence being caused by gas-rich major mergers. ",
        "conclusions": ""
    },
    "0910/0910.2078_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sec:intro}  PHL~932 was first noted as a 12th-magnitude blue star during a survey by Haro \\& Luyten (1962). It was studied in more detail by Arp \\& Scargle (1967), who discovered a faint emission nebula surrounding the star which they assumed to be a planetary nebula (PN).  The nebula is asymmetrically placed around the ionizing star (Figure~\\ref{fig:neb}), which has a position of $\\alpha,\\delta$ = 00$^{h}$59$^{m}$56.67$^{s}$, +15\\arcdeg 44$^{m}$ 13.7$^{s}$ (J2000).  A spectrum of the nebula was obtained by Arp \\& Scargle (1967), who noted strong [O\\,\\textsc{ii}] $\\lambda$3727 and weak H$\\alpha$ emission.  The almost complete absence of [O\\,\\textsc{iii}] emission (see \\S\\ref{sec:em_nebula} below) suggests the nebula has a very low ionization parameter, assuming it is photoionized. An investigation of the stellar spectrum by M\\'endez et al. (1988) placed the central star squarely on the extreme horizontal branch (EHB) and allowed the classification of PHL\\,932 as a hot subdwarf~B (sdB) star.\\footnote{M\\'endez et al. (1988) classified the star as ``late sdO'' (or sdOB) based on their detection of HeII $\\lambda$4686 in absorption.}  Subdwarf B stars are unlike the typical central stars of PN. They do not evolve to the asymptotic giant branch (AGB), but rather move directly on to the white dwarf cooling curve. Their progenitors are solar-like stars that have somehow shed nearly all of their hydrogen envelopes prior to the onset of core helium fusion (e.g. Heber 1986). Exactly how this occurs is uncertain, although it is possible that some kind of binary interaction is involved (Han et al. 2003). Indeed, De Marco et al. (2004) and Af\\c{s}ar \\& Bond (2005) have claimed the radial velocity of PHL\\,932 is variable, which may indicate binarity (however, see Section~\\ref{sec:binary}). More recently, De Marco (2009) has strongly argued for the general importance of binary interactions for the formation and shaping of PN.  Despite the unusual nature of this system, the nebula around PHL~932 has not been studied in much detail since its discovery. In this paper, we present an in-depth analysis of the nebula and the ionizing star based on new and archival multi-wavelength observations. We use the results of this investigation to examine the evolutionary status of the star and to reclassify the nebula.  \\begin{figure}[h] \\begin{center} \\includegraphics[width=7cm]{PHL932BRcomb.ps} \\caption{The asymmetric emission nebula around PHL\\,932.  The star is at the centre of this combined SuperCOSMOS broadband red and blue image, which is 10\\arcmin\\ across.}\\label{fig:neb} \\end{center} \\end{figure} ",
        "conclusions": "\\label{sec:disc}  We have used a combination of new and archival multi-wavelength observations to show that (a) PHL\\,932 is unlikely to be in a close binary system as previously suggested, and (b) the emission nebulosity surrounding PHL\\,932 is \\emph{not} a PN as has been commonly assumed, nor is it physically associated with the star.  We conclude instead that PHL~932 is ionizing a ``clumpy'' region of the ISM as it travels through it; in other words, the emission nebula is simply a Str\\\"omgren zone (or HII region), rather than a PN.  We note that another putative PN, EGB~5 (Ellis, Grayson \\& Bond 1984) is also associated with an sdB star.   An unpublished MSSSO 2.3-m long-slit spectrum of the nebula taken by one of us (D.J.F) shows a nebula of low excitation, similar to that around PHL~932. Like PHL~932, the morphology of EGB 5 on DSS images militates against it being a bona fide PN, though deep narrowband \\ha\\ images are required for a definitive conclusion. We likewise consider the nebula around EGB~5 to be another candidate Str\\\"omgren zone in the ISM.   Such nebulae are seen around hot subdwarfs and white dwarfs when the ISM is sufficiently dense to give an emission measure that facilitates optical detection (cf. Tat \\& Terzian 1999).  In future papers in this series, we will show that several other nearby emission nebulae currently assumed to be PN are also Str\\\"omgren zones in the ISM; see Frew \\& Parker (2006) and Madsen et al. (2006) for preliminary results."
    },
    "0910/0910.5372_arXiv.txt": {
        "abstract": "We present an analysis of the evolution of 8625 Luminous Red Galaxies between $z=0.4$ and $z=0.8$ in the 2dF and SDSS LRG and QSO (2SLAQ) survey. The LRGs are split into redshift bins and the evolution of both the luminosity and stellar mass function with redshift is considered and compared to the assumptions of a passive evolution scenario. We draw attention to several sources of systematic error that could bias the evolutionary predictions made in this paper. While the inferred evolution is found to be relatively unaffected by the exact choice of spectral evolution model used to compute K+e corrections, we conclude that photometric errors could be a source of significant bias in colour-selected samples such as this, in particular when using parametric maximum likelihood based estimators. We find that the evolution of the most massive LRGs is consistent with the assumptions of passive evolution and that the stellar mass assembly of the LRGs is largely complete by $z\\sim0.8$. Our findings suggest that massive galaxies with stellar masses above $10^{11}$M$_\\odot$ must have undergone merging and star formation processes at a very early stage ($z\\gtrsim 1$). This supports the emerging picture of \\textit{downsizing} in both the star formation as well as the mass assembly of early type galaxies. Given that our spectroscopic sample covers an unprecedentedly large volume and probes the most massive end of the galaxy mass function, we find that these observational results present a significant challenge for many current models of galaxy formation. ",
        "introduction": "Luminous Red Galaxies (LRGs) are arguably some of the brightest galaxies in our Universe allowing us to study their evolution out to much higher redshifts than is possible with samples of more typical galaxies. LRGs are known to form a spectroscopically homogenous population that can be reliably identified photometrically by means of simple colour selections \\citep{Eisenstein:01,Cannon:2SLAQ}. The presence of a strong 4000 \\AA break in these galaxies also means that accurate photometric redshifts can be derived for large samples where spectroscopy is difficult to obtain \\citep{Padmanabhan:05, Collister:MegaZ, Abdalla:MegaZ}. Consequently, samples of luminous red galaxies have emerged as an important data set both for studies of cosmology \\citep{Blake:LRG07,Blake:LRG08,Cabre:LRG09} and galaxy formation and evolution \\citep{Wake:06,Cool:08}.  The study of massive galaxies in our Universe is particularly interesting since these galaxies present a long-standing problem for models of galaxy formation. While in the standard $\\Lambda$CDM cosmology, massive galaxies are thought to be built up through successive mergers of smaller systems \\citep{WhiteRees:78}, many observations now suggest that these galaxies were already well assembled at high redshifts \\citep{Glazebrook:04, Cimatti:06, Scarlata:07, Ferreras:09, Pozzetti:09}. At the same time, star formation and merging activity has also been seen in such systems between $z\\sim1$ and the present day \\citep{Lin:04, Stanford:04} and the studies of \\citet{Bell:04} and \\citet{Faber:07} have found evidence that the luminosity density of massive red galaxies has remained roughly constant since z$\\sim$1 implying a build up in the number density of these objects through mergers. Most observational results now support downsizing in star formation - i.e most massive galaxies having lower specific star formation rates than their less massive counterparts - and many galaxy formation models are able to reproduce these observations e.g. by invoking quenching mechanisms such as AGN feedback \\citep{Croton:06, DeLucia:06}. However, the more recently observed trend of downsizing in the mass assembly \\citep{Cimatti:06, Pozzetti:09} presents more of a challenge for galaxy formation models which typically predict that the most massive galaxies were assembled later than their less massive counterparts.  Many of the contradictions in observational studies of massive galaxy formation arise principally for two reasons. Firstly, the way in which early-type galaxies are selected in different surveys can be considerably different. Morphological selection compared to colour selection of such objects will almost certainly result in different samples being chosen, particularly at high redshifts. The colour selection is also usually different for different samples of massive galaxies as we will point out later in this paper. Secondly, many of the studies of massive galaxy formation mentioned so far are deep and narrow spectroscopic samples that will suffer from biases due to the effects of cosmic variance. In addition, one has to be careful in interpreting observational results as in current galaxy formation models, star formation and mass assembly in galaxies are not necessarily concomitant. So while the stars in early-type galaxies may have formed very early in the Universe's history, they may have formed in relatively small units and only merged at lower redshifts to create the massive ETGs we see today.  Studies of the evolution of luminous red galaxies such as those analysed in this paper have already supported the evidence that massive galaxies have been passively fading and show very little recent star formation \\citep{Wake:06, Cool:08}. In this paper, we extend this work by also considering the mass assembly of these systems. The 2dF and SDSS LRG and QSO (2SLAQ) survey presents a comprehensive improvement in volume for massive galaxy samples. The survey covers an area of 186 deg$^2$ and reaches out to a redshift of 0.8 making it a promising data set to study the evolution of massive red galaxies. Furthermore, much of the 2SLAQ area is now covered by the UKIDSS Large Area Survey (LAS) that provides complementary data in the near infra-red bands for the optically selected LRGs. The advantage of near infra-red data is that the mass-to-light ratios and k-corrections are largely insensitive to the galaxy or stellar type and therefore the total infra-red flux in for example the K-band provides a good estimate of the total stellar mass of the galaxies. This stellar mass estimate allows us to study not only the star formation history but also the mass assembly history of these systems. As is the case for optically selected spectroscopic samples of massive galaxies, the K-band selected surveys have so far been restricted to relatively small and deep patches of the sky \\citep{Mignoli:K20, Conselice:07}. Clearly a large spectroscopic survey of massive galaxies with optical and near infra-red photometry, will allow for better constraints on the evolution of these systems.  In this paper, we utilise a spectroscopic sample of colour-selected massive red galaxies from the 2SLAQ survey between redshifts 0.4 and 0.8. Optical photometry is obtained from the Sloan Digital Sky Survey supplemented with near infra-red data from the UKIDSS Large Area Survey (LAS). We consider the evolution of these galaxies in terms of their observed colours, luminosities and comoving number densities focussing particularly on the effects of colour selection on the inferred evolution.  \\citet{Wake:06} have presented a comprehensive analysis of the luminosity function of Luminous Red Galaxies using data from both the Sloan Digital Sky Survey and the 2SLAQ Survey. These authors use a subset of data from the 2SLAQ survey in the redshift range 0.5 to 0.6 and by comparing this galaxy population to a lower redshift population at $0.17<z<0.24$ from SDSS, they are able to establish that the LRG LF does not evolve beyond that expected from a simple passive evolution model at these redshifts. Meanwhile, \\citet{Cool:08} have compared the low-redshift SDSS LRG population to their own high-redshift sample at redshifts of $\\sim 0.9$ and once again find little evidence for evolution beyond the passive fading of the stellar populations. We extend the work of these authors by considering now most of the galaxies available in the 2SLAQ data set. This enables the redshift range of 2SLAQ galaxies used to be broadened to $0.4\\leq z < 0.8$.  The 2SLAQ luminosity function is presented for 8625 LRGs as compared to the 1725 used by \\citet{Wake:06}. These results are useful in filling the gap between redshift 0.4 and 0.5 and redshift 0.6 and 0.8 in the \\citet{Wake:06} and \\citet{Cool:08} data sets. The sample is split into four redshift bins and the evolution of the luminosity and colour of LRGs between these redshift bins is considered. In addition, we also consider the stellar mass function and its evolution with redshift. Furthermore, we use various different stellar population synthesis models that are commonly used in the literature, to model the LRGs, including the new models of \\citet{Maraston:09} and quantify the sensitivity of the luminosity function estimate to changes in these models. The optical colours of LRGs have long been difficult to model using standard spectral evolution models \\citep{Eisenstein:01} and this problem has only recently been solved \\citep{Maraston:09} thereby allowing us to utilise the most accurate spectral evolution models of LRGs to date in order to infer their evolution.  The paper is structured as follows. In $\\S$ \\ref{sec:data} we describe the spectroscopic 2SLAQ data set as well as SDSS and UKIDSS photometry for these galaxies. $\\S$ \\ref{sec:models} describes the spectral evolution models used in this paper to model the LRGs. We consider the optical luminosity function of the 2SLAQ LRGs in $\\S$ \\ref{sec:lf} and its evolution with redshift as well as its sensitivity to cosmic variance, photometric errors and changes in the spectral evolution models. In $\\S$ \\ref{sec:number} we present estimates for the LRG mass function in different redshift bins and analyse its sensitivity to the choice of spectral evolution model as well as the IMF. Finally, we discuss our results in $\\S$ \\ref{sec:disc} in terms of models of galaxy formation and evolution. Throughout this paper we assume a cosmological model with $\\Omega_m$=0.3, $\\Omega_\\Lambda$=0.7 and $h$=0.7. All magnitudes are in the AB system unless otherwise stated. ",
        "conclusions": "This paper has examined the evolution of the optical and near infra-red luminosity function as well as the stellar mass function of 8625 LRGs from the 2SLAQ survey between redshift 0.4 and 0.8. We have demonstrated the unique position of 2SLAQ LRGs among other massive galaxy samples due to the large volume probed by the sample as well the availability of spectroscopic redshifts for all galaxies and near infra-red data for $\\sim$75\\% of them. This has allowed us to probe the very high-mass end of the stellar mass function which places the most stringent constraints on models of galaxy formation. We have also used the spectral evolution models of \\citet{Maraston:09} which are the most accurate models of LRGs to date to study the evolution of these objects although our conclusions change little if other models are used instead. Specifically we draw the following conclusions:  \\begin{itemize}  \\item{The evolution of the optical luminosity function of LRGs between redshifts 0.4 and 0.8 is consistent with the assumptions of a passive evolution model where the stars were formed at high redshift with little or no evidence for recent episodes of star formation.}  \\item{A downturn is seen at the faint end of the LRG luminosity function due to the effects of downsizing in the star formation. The faintest galaxies are expected to be less massive and therefore bluer than the bright objects and such objects are excluded from the LRG sample due to the imposed colour selection of these objects which selects only the reddest galaxies.}  \\item{We find that photometric errors may induce a significant bias into our sample of LRGs and scatter galaxies both in and out of our sample across the colour selection boundaries. We show that more galaxies are likely to be scattered into the sample than are scattered out leading to an overprediction of the observed space density but the shape of the luminosity function does not change if we remove the galaxies with large photometric errors.}  \\item{The LRG luminosity function is found to be unaffected by cosmic variance due to the large volume occupied by the sample.}  \\item{The luminosity function is also relatively insensitive to the choice of spectral evolution model and the metallicity although models with some residual star formation shift the luminosity function towards fainter absolute magnitudes.}  \\item{The stellar mass function for these LRGs for $M > 3 \\times 10^{11} M_\\odot$ also shows little evidence for evolution between redshifts 0.4 and 0.8 suggesting that these most massive systems were already well assembled at redshifts of 0.8. This is consistent with the emerging picture of downsizing in the mass assembly of massive galaxies.}  \\item{The stellar mass function estimate is also relatively insensitive to the choice of spectral evolution model assumed in the calculation of $V_{max}$. However, different choices of the stellar initial mass function will shift the mass function estimate as found in previous studies.}  \\item{The comoving number density of LRGs with $M> 3 \\times 10^{11} M_\\odot$ has changed little between redshifts 0.8 and 0.4.  The same is true for LRGs with $M > 10^{12} M_\\odot$. This is consistent with other observational results for massive galaxy samples and does not agree with the predictions of most current galaxy formation models. We find that the models of \\citet{Granato:04} may be promising in matching our observations although they are yet to be compared with massive galaxy samples at intermediate redshifts.}  \\end{itemize}  Overall, our results support the emerging picture of downsizing in both the star formation and mass assembly of early type galaxies and present a significant challenge for current models of galaxy formation. We have shown our findings to be robust to changes in spectral evolution models, cosmic variance and photometric errors and conclude that these new observations of the most massive and most luminous galaxies in our Universe will need to be reconciled with the models if progress is to be made in the field of massive galaxy formation and evolution."
    },
    "0910/0910.2990_arXiv.txt": {
        "abstract": "We present new near-IR \\hh, CO J=2-1, and CO J = 3-2 observations to study outflows in the massive star forming region IRAS 05358+3543. The Canada-France-Hawaii Telescope \\htwo\\ images and James Clerk Maxwell Telescope CO data cubes of the IRAS 05358 region reveal several new outflows, most of which emerge from the dense cluster of sub-mm cores associated with the \\necluster\\ cluster to the northeast of \\region. We used Apache Point Observatory (APO) JHK spectra to determine line of sight velocities of the outflowing material. Analysis of archival VLA cm continuum data and previously published VLBI observations reveal a massive star binary as a probable source of one or two of the outflows.  We have identified probable sources for 6 outflows and candidate counterflows for 7 out of a total of 11 seen to be originating from the \\region\\ clusters.  We classify the clumps within \\necluster\\ as an early protocluster and \\swcluster\\ as a young cluster, and conclude that the outflow energy injection rate approximately matches the turbulent decay rate in \\necluster. ",
        "introduction": "Collimated, bipolar outflows accompany the birth of young stars from the earliest stages of star formation to the end of their accretion phase \\citep[e.g.][]{reipurth2001}.    While the birth of isolated low-mass stars is becoming well understood, the formation of massive stars ($>10 \\msun$) and clusters remains a topic of intense study.    Observations show that moderate to high-mass stars tend to form in dense clusters \\citep{lada2003}.    In a clustered environment,  the dynamics of the gas and stars can profoundly impact both accretion and mass-loss processes. Feedback from these massive clusters may play a significant role in momentum injection and turbulence driving in the interstellar medium.  Outflows from massive stars are less studied than those from low mass stars largely because massive stars accrete most of their mass while deeply embedded. Therefore, unlike low mass young stars that are accessible in the optical, massive stellar outflows can only be seen at infrared and longer wavelengths. Direct evidence for jets from massive young stellar objects (YSOs) from \\hh\\ or optical emission is generally lacking \\citep[e.g.][]{alvarez2005,kumar2002,wang2003}, although there is evidence that massive stars are the sources of collimated molecular outflows from millimeter observations \\citep[e.g.][]{beuther2002b}.  Outflows from massive stars may allow accretion to continue after their radiation pressure would otherwise halt accretion in a spherically symmetric system \\citep{krumholz2009}.  They therefore represent a crucial component in understanding how stars above $\\sim$10 \\msun\\ can form.  \\region\\ is a double cluster of embedded infrared sources located at a distance of 1.8 kpc in the Auriga molecular cloud complex \\citep{heyer1996} associated with the HII regions Sh-2 231 through 235 at Galactic coordinates around $l,~b$ = 173.48,+2.45 in the Perseus arm.   \\necluster\\ is the collection of highly obscured and mm-bright sources slightly northeast of \\swcluster, which is the location of the IRAS 05358+3543 point source and the optical emission nebula (see Figure \\ref{fig:overview_ha}).  The IRAS source is probably a blend of the three brightest infrared objects in the MSX A-band and MIPS 24 \\um\\ images, which are located at \\necluster, IR 41, and IR 6. For the purpose of this paper,the whole complex including both sources is referred toas \\region, and otherwise refer to individual objects specifically.  Early observations revealed the presence of OH \\citep{Wouterloot1993}, \\HtwoO\\ \\citep{Scalise1989, Henning1992}, and methanol \\citep{Menten1991} masers about an arcminute northeast of the IRAS source, indicating that massive stars are likely present at that location.  Near infrared observations revealed the presence of two embedded clusters  \\citep{porras2000,jiang2001} labeled \\swcluster\\ for the southwestern cluster associated with the IRAS source, and \\necluster\\ for the northeastern cluster located near the OH, \\HtwoO, and CH$_3$OH masers. Stars identified in \\citet{porras2000} are referred to by the designation ``IR (number)'' corresponding to the catalog number in that paper. \\citet{porras2000} also included scanning Fabry-Perot velocity measurements of the inner $\\sim1$\\arcmin.  CO observations revealed broad line wings indicative of a molecular outflow \\citep{casoli1986,shepherd1996}.  \\citet{kumar2002} and \\citet{khanzadyan2004} presented narrow band images of 2.12 \\um\\ \\htwo\\ emission that reveled the presence of multiple outflows.  Interferometric imaging of CO and SiO confirmed the presence of at least three flows emerging from the northeast cluster centered on the masers \\citep{beuther2002} having a total mass of about 20 \\msun .  \\citet{beuther2002} also presented MAMBO 1.2 mm maps and a mass estimate of 610 \\msun\\ for the whole region. \\citet{williams2004} presented SCUBA maps and mass estimates of the clusters of 195/126\\msun\\ for \\necluster\\ and 24/12 \\msun\\ for \\swcluster\\ (850 \\um/450 \\um). \\citet{Zinchenko1997} measured the dense gas properties using the \\ammonia\\ (1,1) and (2,2) lines.  They measure a mean density $n \\approx 10^{3.60}$ \\percc, temperature 26.5K, and a mass of 600 \\msun .  The total luminosity of the two clusters is about 6300 \\lsun , indicating that the region is giving birth to massive stars \\citep{porras2000}.  Millimeter wavelength interferometry with arcsecond angular resolution has revealed a compact cluster of deeply embedded sources centered on the \\HtwoO\\ and methanol maser position \\citep{beuther2002,beuther2007,leurini2007}. \\citet{beuther2002} identified 3 mm continuum cores, labeled mm1-mm3 (shown in Figure \\ref{fig:outflowsh2}).  \\citet{beuther2007} resolved these cores into smaller objects.  Source mm1a is associated with a cm continuum point source and will be discussed in detail below.  \\region\\ has previously been observed at low spatial resolution in the J=2-1 and J=3-2 transitions with the Kosma 3m telescope \\citep{Mao2004}.  While the general presence of outflows was recognized and a total mass estimated, the specific outflows were not resolved.  \\citet{beuther2002} observed the CO J=6-5, J=2-1, and J=1-0 transitions at moderate resolution in the inner few arcminutes. \\citet{Thomas2008} observed C$^{17}$O in the J=2-1 and J=3-2 transitions with a single pointing using the JCMT. ",
        "conclusions": "We have presented a multiwavelength study of the \\region\\ star forming region. \\region\\ contains an embedded cluster of massive stars and is surrounded by outflows.  The outflows were linked to probable sources and determined that at least one outflow is probably associated with a massive ($\\sim 10 \\msun$) star. Added kinematic information and a wide field view of the infrared outflows has been used to develop a more complete picture of the region.  \\begin{itemize} \\item \\necluster\\ is a Protocluster and \\swcluster\\ is a Young Cluster \\item Energy injection on the scales of \\region\\ can maintain turbulence, but on the small scales of the \\necluster\\ protocluster, is inadequate by $\\sim 2$ orders of magnitude.  \\necluster\\ is collapsing. \\item there are 11 candidate outflows, 7 of which have candidate counterflows, in the \\region\\ complex \\item there is a probable massive binary with one member of mass 12 \\msun\\ in mm1a, and the other which is the source of Outflow 2 \\item there are at least two moderate-mass ($\\sim$5\\msun) young stars in \\region\\ \\end{itemize}  We have identified additional middle- and high-mass young stars with outflows, and presented a case for a high-mass binary system within the millimeter core mm1a."
    },
    "0910/0910.5054_arXiv.txt": {
        "abstract": "{ Cosmology contributes a good deal to the investigation of variation of fundamental physical constants. High resolution data is available and allows for detailed analysis over cosmological distances and a multitude of methods were developed. The raised demand for precision requires a deep understanding of the limiting errors involved. The achievable accuracy is under debate and current observing proposals max out the capabilities of todays technology. The question for self-consistency in data analysis and effective techniques to handle unknown systematic errors is of increasing importance. An analysis based on independent data sets is put forward and alternative approaches for some of the steps involved are introduced. ",
        "introduction": "This work is motivated by numerous findings of different groups that partially are in disagreement witch each other. A large part of these discrepancies reflect the different methods of handling systematic errors. Evidently systematics are not yet under control or fully understood. We try to emphasize the importance to take these errors, namely i.e. calibration issues, into account and put forward some measures adapted to the problem. ",
        "conclusions": ""
    },
    "0910/0910.2397_arXiv.txt": {
        "abstract": "In this paper bulk viscosity is introduced to describe the effects of cosmic non-perfect fluid on the cosmos evolution and to build the unified dark energy (DE) with (dark) matter models. Also we derive a general relation between the bulk viscosity form and Hubble parameter that can provide a procedure for the viscosity DE model building. Especially, a redshift dependent viscosity parameter $\\zeta\\propto\\lambda_{0}+\\lambda_{1}(1+z)^{n}$ proposed in the previous work by X.H.Meng and X.Dou in 2009\\cite{md} is investigated extensively in this present work. Further more we use the recently released supernova dataset (the Constitution dataset) to constrain the model parameters. In order to differentiate the proposed concrete dark energy models from the well known $\\Lambda$CDM model, statefinder diagnostic method is applied to this bulk viscosity model, as a complementary to the $Om$ parameter diagnostic and the deceleration parameter analysis performed by us before. The DE model evolution behavior and tendency are shown in the plane of the statefinder diagnostic parameter pair \\{$r,s$\\} where the fixed point represents the $\\Lambda$CDM model. The possible singularity property in this bulk viscosity cosmology is also discussed to which we can conclude that in the different parameter regions chosen properly, this concrete viscosity DE model can have various late evolution behaviors and the late time singularity could be avoided. We also calculate the cosmic entropy in the bulk viscosity dark energy frame, and find that the total entropy in the viscosity DE model increases monotonously with respect to the scale factor evolution, thus this monotonous increasing property can indicate an arrow of time in the universe evolution, though the quantum version of the arrow of time is still puzzling. ",
        "introduction": "Type Ia supernova and other astrophysics observations together indicate that our universe is accelerating now \\cite{sn1}. Different models are proposed to try to describe or understand this surprisingly exotic phenomenon. If we make the assumption that general relativity is still correct to the scale of cosmos, an effective term contributes to negative pressure should be added to the right hand side of Einstein's field equation in the general theory of relativity to explain the recent stage speed-up of our observational universe expansion. This is the basic idea to the so called dark energy concept. So far for the ten years' old DE phenomena there are many both theoretical and observational attempts to understand the mechanism. The introduction of the cosmological constant corresponds to a negative pressure fluid with a specially constant density all the time which has been playing a key role in the universe evolution by uniformly distributed over the whole cosmic space-time and media, but the cosmological constant existence raises serious fundamental physics problem, the so-called cosmological constant problem for both new (coincidence) and old (so tiny) ones. Another class of models try to modify the traditional Einstein's general theory of relativity to the large cosmology scale, by arguing that the recently appearing acceleration phase of the unverse expansion comes from the break down of general relativity in cosmic scale. The so-called $f(R)$ gravity, to mention one for example, which generalizes the Hilbert-Einstein action is categorized in this class \\cite{mw1}. (For more details and references, you may see the recent reviews \\cite{r1} \\cite{r2}).  In the context of perfect fluid, some models based on fluid mechanics method are studied extensively, such as the Chaplygin gas and generalized Chaplygin gas models which modify the equation of state, and barotropic fluids dark Energy \\cite{ls}. Also for the purpose to consider more realistic situation in the evolution of the universe, the concept of viscosity is introduced into the investigation of the cosmos evolution from fluid mechanics \\cite{v2} \\cite{v3} \\cite{v7} \\cite{v8} \\cite{v9} \\cite{v10} \\cite{v11}. Earlier attempt \\cite{pc} in this area even ``predicts'' the late acceleration of the universe expansion. The dissipative effect in the fluid is always due to shear and bulk viscosity characterized by shear viscosity parameter $\\eta$ and bulk viscosity parameter $\\zeta$. In the study of cosmology, the shear viscosity disappears in the Friedmann-Robertson-Walker's isotropic and homogeneous framework with the largest spherical symmetry. Only bulk viscosity could play a role in the realistic models. Its effects can be shown from an added correction term with the minus sign to cosmic pressure $\\tilde p=p-3\\zeta H$. This composite formula motives the trial on the connection between bulk viscosity of the cosmic fluid and dark components (energy and matter). As we understand so far that purely gravitational probes can only provide information on a single effective matter-energy fluid, which may consist of the dark energy and (dark) matter composites as well as the least dominated radiation energy at present universe, so for the dominated two functioning dark components we can describe them as a single dark fluid \\cite{Ren}. In this paper, we discuss a concrete unified model building by paying attention on the modification of the cosmic media from the simple assumption as a perfect fluid to a piratical viscous fluid encoded in the energy-stress tensor contents, that is, at the right hand side of the Einstein's field equation. For the bulk viscosity DE model building a remark is needed here. Considering the observational dark energy ingredient fraction today, it takes about $\\frac{2}{3}$ of the whole matter-energy density, so the corresponding pressure provided by the bulk viscosity correction dominantly surpasses the remaining pressure contributions from other cosmic energy-matters. This is obviously different from the traditional non-equilibrium thermodynamics, in which the viscosity contribution is only a little correction to the pressure term. Some researches try to find a mechanism to support such a fluid behavior by a kind of non-standard interaction introduced between the dark matter and energy components \\cite{inter1} \\cite{inter2}. It is certainly important for these attempts to find more support or hints behind cosmic observations in the study of bulk viscosity cosmology along this possible line.  For concrete model building, detailed form of viscosity parameter is needed to settle down the theoretical framework. In the reference \\cite{Ren}, a scale factor dependent viscosity is proposed, which is different from the only density dependent form \\begin{equation} \\zeta=\\zeta_{0}+\\zeta_{1}\\frac{\\dot a}{a}+\\zeta_{2}\\frac{\\ddot a}{\\dot a} \\end{equation} It could be shown that it is equivalent to a modified equation of state(EOS) \\begin{equation} p=(\\gamma-1)\\rho+p_{0}+w_{H}H+w_{H2}H^{2}+w_{dH}\\dot H. \\end{equation} And other interesting physical properties have been investigated in that kind of models \\cite{rm} \\cite{con}. In this present paper, by largely extending its contents we will continue the study of a redshift dependent viscosity form $\\zeta\\propto\\lambda_{0}+\\lambda_{1}(1+z)^{n}$ as proposed in the previous work \\cite{md}.  We concentrate on the late time singularity discussion and its entropy expression of the bulk viscosity DE model. For the phantom dark energy model, there exists a cosmic singularity in the future cosmos evolution, that is, the so-called cosmic ``doomsday''. As shown below, we could see that this viscosity DE model represents different evolution behaviors, and the future cosmic singularity will disappear under some proper parameter ranges selection. From this view of point, the model has the flexibility to produce either quintessence or phantom properties, that is, we can easily achieve its EoS either larger or less than characteristic -1.  An important tool for investigating dark energy model characters nowadays is by the introduction of some geometry quantities, for instance the statefinder diagnostic parameters \\cite{sf}, which are quantities dependent on high order derivatives of the scale factor, such as to the $\\dddot a$. The usual statefinder parameter pair \\{$r,s$\\} of the model concerned is calculated explicitly to demonstrate the viscosity DE model behaviors. In our previous work, deceleration parameter and $Om$ diagnostic parameter \\cite{om} are performed to show the properties. We have found that in statefinder pair plane, our model and $\\Lambda$CDM model could be easily discriminated in some redshift ranges. We prospect that increasing the quality and quantity of measurement data will enhance our ability to make accurate discrimination of different current cosmology models and rule out some. At the same time, new diagnostic methods merit further investigations, especially diagnostic quantity in the higher order perturbation level.  The paper is organized as follows: in the second section, some general features and remarks about viscosity dark energy model are given, and we further discuss the redshift dependent model. In the third section, we give the singularity discussion of our model. Some solutions are given there. In Sec. \\textbf{IV}, we calculate entropy of this viscosity model. Finally, conclusion and discussions are presented. We leave the data fitting procedure in the appendix. ",
        "conclusions": "In this letter we continue and largely extended our previous work on the single and unified bulk viscous fluid as a potential dark energy candidate by presenting an explicit viscosity form to mimic dark energy behaviors and confront it with current observational data sets. The specific feature here is a variable coefficient for the new bulk viscosity form proposed, characterized by two free parameters that can be best fitted by astrophysics observational data sets. the best fitting results have shown that this concrete model could yield theoretical prediction values in an acceptable level by working out the numerical processing to the latest released joint observational data sets.  Furthermore, we have performed the statefinder diagnostic parameter analysis to this unified viscosity DE model, finding that in most evolution stages of the unverse, statefinder parameters could be used to obviously distinguish this viscosity DE model from the $\\Lambda$CDM model. But as shown in the figure of the statefinder pair parameter \\{$r,s$\\}, the viscosity DE model evolution trajectory passes the special point corresponding to the well known $\\Lambda$CDM model, where the models degeneration emerges. The new $Om$ parameter diagnostics made previously could be a more powerful tool then, which might discriminate concrete models from the $\\Lambda$CDM model in the whole evolution history.  In this present work, we particularly concentrate on the singularity behavior of the unified viscosity dark energy model. We find that different parameter range selection, especially the region of power parameter $n$, influences the finite future singularity. For the $n=0$ case, which corresponds to the case with a constant viscosity, there is an exact solution for the evolution scale factor. For the early universe and the $n<0$ case, the solution of the scale factor $a(t)$ has possessed the same structure. A further study of viscosity effects on the early universe evolution, especially during inflation stage, seems very interesting with rich possibilities.  In the context of the unified viscosity DE cosmology, we also calculate the entropy for the total evolution universe, which has been expressed in a complex form with this non-perfect viscosity media. Calculations of the total entropy for the viscosity universe evolution represent that the general second law of thermodynamics holds in the whole or global universe description. The worked out results show that in this complicated context, the total entropy is increasing with the redshift decreasing, or cosmic time flying, including the cosmic expansion accelerating stage as observational data indicating now, which has been plotted in figure 5 with three free parameter values chosen. The monotonous increasing property of the total entropy with cosmic time flowing, the viscosity DE universe evolution may provide us an arrow of time to describe the complex universe changing direction. Though a definite concept for the arrow of time is debating, now very puzzling in different situations or versions: classical physics, quantum physics, cosmology and quantum gravity, especially it is essential for a correct quantum gravity theory to be expected to appear, the thermodynamics second law is generally believed to hold to describe the global evolution of our observational universe. So the reasonable entropy expression is richly encoded helpful information for the concerned system. It is certainly interesting and worthy of further efforts.  Dark energy physics involves many fundamental concepts and beyond in our already established \"standard models\" for both particle physics and gravity. Viscosity media seems to relate the matter sector to the geometric gravity side via the Einstein's general relativity equation. Actually it also can be reconstructed effectively from the left side to the right side of the equation by modified gravity or extra dimension models for example, too, which is also intriguing. With the available and upcoming high quality and increasingly large amount of astrophysics observational data, especially the good low redshift SNe Ia data sets with less uncertainty for possible errors from the dust effects alleviated under control we expect the ten-years old mysterious dark energy problem will be pinned down not too long compared with the long time standing unsolved dark matter mystery. Maybe the practically unified viscosity DE model or the like can provide an economic mechanism to answer the both uncovered secrets in one dark sector, a tale for the two mysteries."
    },
    "0910/0910.3672_arXiv.txt": {
        "abstract": "{We present results of variability search in the field of the young open cluster NGC\\,1502. Eight variable stars were discovered. Of six other stars in the observed field that were suspected for variability, we confirm variability of two, including one $\\beta$~Cep star, NGC\\,1502-26.  The remaining four suspects were found to be constant in our photometry. In addition, $UBVI_{\\rm C}$ photometry of the well-known massive eclipsing binary SZ~Cam was obtained.  The new variable stars include: two eclipsing binaries of which one is a relatively bright detached system with an EA-type light curve, an $\\alpha^2$\\,CVn-type variable, an SPB candidate, a field RR Lyr star and three other variables showing variability of unknown origin. The variability of two of them is probably related to their emission in H$\\alpha$, which has been measured by means of the $\\alpha$ index obtained for 57 stars brighter than $V \\approx$16 mag in the central part of the observed field.  Four other non-variable stars with emission in H$\\alpha$ were also found.  Additionally, we provide $VI_{\\rm C}$ photometry for stars down to $V =$ 17~mag and $UB$ photometry for about 50 brightest stars in the observed field. We also show that the 10-Myr isochrone fits very well the observed color-magnitude diagram if a distance of 1~kpc and mean reddening, $E(V-I_{\\rm C}) =$ 0.9~mag, are adopted.} {stars: early-type -- stars: emission-line, Be --- open clusters and associations: individual: NGC\\,1502} ",
        "introduction": "This is the seventh paper in the series in which we present results of the variability search for B-type variables (mainly $\\beta$~Cep, SPB and massive binaries) in the young open clusters of the Northern Hemisphere. The results for previously studied clusters were briefly summarized by Pigulski \\etal (2002).  In addition to the discovery or verification of variability of about 70 stars, mostly members of the investigated clusters, it appeared that the incidence of $\\beta$~Cep stars in the northern clusters is lower than in the southern ones. This was interpreted in terms of the metallicity gradient in the Galaxy. One of the main motivations for the search is also a selection of clusters containing large numbers of pulsators suitable for asteroseismology.  An example is the previous paper of the series (Ko\\l{}aczkowski \\etal 2004), in which we investigated the open cluster NGC\\,6910 discovering four $\\beta$~Cep stars. This prompted us to choose NGC\\,6910 as a target of a two-year photometric and spectroscopic campaign.  The data from the campaign are in the course of reduction, but preliminary analysis already increased the number of $\\beta$~Cep stars in this cluster to seven (Pigulski \\etal 2007) making it an excellent target for seismic modelling and one of very few which are rich in $\\beta$~Cep stars.  In the present paper, we report the results of the variability search in the field of another young open cluster, NGC\\,1502. For the first time, the data from the new CCD camera, with a much larger field of view, are included. ",
        "conclusions": "The open cluster NGC\\,1502 contains only a single low-amplitude $\\beta$~Cep star (NGC\\,1502-26). This appears to be in a contrast to NGC\\,6910 (Ko{\\l}aczkowski \\etal 2004, Pigulski \\etal 2007) which is not very rich in stars either but harbors as many as seven $\\beta$~Cep stars.  Whether the difference can be interpreted in terms of a difference in metallicity is not known. Unfortunately, the only metallicity determination that can be related to NGC\\,1502 is that from Str\\\"omgren photometry by Eggen (1985) for HD\\,25056 in Cam~OB1. He reported [Fe/H] $=$ 0.16~dex for this star.  The value of [Fe/H] is rather uncertain. In addition, the star might not be related to NGC\\,1502 at all. Definitely, a spectroscopic determination of abundances of stars in the cluster are highly desirable.  SZ Cam remains one of the most interesting stars of the cluster being an excellent case of a massive hierarchical system that are found frequently in the cores of open clusters.  Thorough studies of such systems may shed light on the role of binaries in massive star formation in general. SZ~Cam was also succesfully used to derive the distance to the cluster (Gorda \\etal 2007) and we may hope that further observations will improve this determination.  The relatively bright EA-type system (NGC\\,1502-7) we discovered might be also used for this purpose.  \\Acknow{This research has made use of the WEBDA database, operated at the Institute for Astronomy of the University of Vienna. This work was supported by two grants: N\\,N203\\,302635 from Polish MNiSzW and Chilean Proyecto FONDECYT Nr 3085010. We thank Prof.~M.\\,Jerzykiewicz for his comments made upon reading the manuscript. We also thank J.\\,Molenda-\\.Zakowicz, G.\\,Kopacki and Z.\\,Ko{\\l}aczkowski for making some observations of the cluster.}"
    },
    "0910/0910.4863_arXiv.txt": {
        "abstract": "The centers of most galaxies in the local universe are occupied by compact, barely resolved sources. Based on their structural properties, position in the fundamental plane, and integrated spectra, these sources clearly have a stellar origin. They are therefore called \"nuclear star clusters\" (NCs) or \"stellar nuclei\". NCs are found in galaxies of all Hubble types, suggesting that their formation is intricately linked to galaxy evolution. Here, I review some recent studies of NCs, describe ideas for their formation and subsequent growth, and touch on their possible evolutionary connection with both supermassive black holes and globular clusters. ",
        "introduction": "The nuclei of galaxies are bound to provide ``special'' physical conditions because they are located at the bottom of the potential well of their host galaxies. This unique location manifests itself in various distinctive phenomena such as super-massive black holes (SMBHs), active galactic nuclei (AGN), central starbursts, or extreme stellar densities. The evolution of galactic nuclei is closely linked to that of their host galaxies, as inferred from a number of global-to-nucleus scaling relations discovered in the last decade.  Recently, observational and theoretical interest has been refocused onto the compact and massive star clusters found in the nuclei of galaxies of all Hubble types. These ``nuclear star clusters'' (NCs) are intriguing objects that are linked to a number of research areas: i) they are a promising environment for the formation of massive black holes because of their extreme stellar density, ii) they may also constitute the progenitors of at least some halo globular clusters via ``NC capture'' following the tidal disruption of a satellite galaxy, and iii) their formation process is influenced by (and important for) the central potential, which in turn governs the secular evolution of their host galaxies.  In what follows, I briefly summarize what has been learned about NCs over the last few years, describe some proposed formation mechanisms of NCs, and discuss the new paradigm of ``central massive objects'' which links NCs with SMBHs in galactic nuclei. Lastly, I briefly mention a scenario in which NCs may be the progenitors of (some) globular clusters. ",
        "conclusions": ""
    },
    "0910/0910.3444_arXiv.txt": {
        "abstract": "Conventionally, long GRBs are thought to be caused by the core collapses of massive stars. During the lifetime of a massive star, a stellar wind bubble environment should be produced. Furthermore, the microphysics shock parameters may vary along with the evolution of the fireball. Here we investigate the variation of the microphysics shock parameters under the condition of wind bubble environment, \\textbf{and allow the microphysics shock parameters to be discontinuous at shocks in the ambient medium. It is found that our model can acceptably reproduce the rebrightenings observed in GRB afterglows, at least in some cases.} The effects of various model parameters on rebrightenings are investigated. The rebrightenings observed in both the R-band and X-ray afterglow light curves of GRB 060206, GRB 070311 and GRB 071010A are reproduced in this model. ",
        "introduction": "Gamma-ray bursts (GRBs) are attractive astrophysics phenomena and had puzzled astronomers for about twenty-four years after their discovery in 1973 (Klebesadel et al. 1973). The discovery of long-lived, multi-band counterparts of GRBs, namely, the afterglows of GRBs, in 1997 is a watershed in GRB research (Costa et al. 1997; Van Paradijs et al. 1997; Frail et al. 1997). Soon after that, the so-called fireball model is recognized as the standard model in view of the fact that it can explain most features of GRB observations well. However, as the advance of observation techniques and the accumulation of observational data, especially after the launch of \\emph{Swift} satellite (Gehrels 2004), a lot of unexpected behaviours appear in GRB afterglows, such as the canonical steep-shallow-normal decay and flares in X-ray afterglows, the flattish decay phase and various rebrightenings in optical afterglows (For review, see Zhang 2007).  In fact, there are more and more rebrightenings detected in GRB afterglows, including GRB 970508 (Galama et al. 1998a), GRB 990123 (Sari \\& Piran 1999), GRB 021004 (Lazzati et al. 2002), GRB 030329 (Berger et al. 2003), GRB 050525A (Blustin et al. 2006), GRB 050820A (Cenko et al. 2006), GRB 050721 (Antonelli et al. 2006), GRB 060206 (Stanek et al. 2007; Liu et al. 2008), GRB 070125 (Updike et al. 2008), GRB 070311 (Guidorzi et al. 2007a), GRB 071003 (Perley et al. 2008), GRB 071010A (Covino et al. 2008). Many different mechanisms have been proposed to explain these rebrightenings, such as density jump (Lazzati et al. 2002; Dai \\& Wu 2003; Tam et al. 2005), energy injection (Huang et al. 2006), two-component jet (Huang et al. 2004, 2006; Liu et al. 2008), reverse shock (Sari \\& Piran 1999), reverberation of the energy input measured by prompt emission (Vestrand et al. 2006), turn-on of the external shock (Stanek et al. 2007; Molinari et al. 2007), large angle emission (Panaitescu \\& Kumar 2007), and spectral peak of existing forward shock (Shao \\& Dai 2005).  Among all the mechanisms, the density jump model needs to be paid special attention to, because density jump in surrounding medium of GRBs is a very natural hypothesis. Since there are more and more examples indicating that some GRBs are associated with Type Ic supernovae (e.g. SN 1998bw/GRB980425, Galama et al. 1998b; SN2003dh/GRB030329, Price et al. 2003) and many host galaxies of GRBs are in process of active star formation (Fruchter et al. 1999; Djorgovski et al. 1998), it is believed that the progenitors of long GRBs are massive Wolf-Rayet (WR) stars (Woosley 1993). Massive stars usually produce very strong stellar wind to push the initial interstellar medium (ISM) in their neighborhood away during their lifetime. The surrounding of these GRBs should be a wind bubble following the density profile $\\rho \\propto r^{-2}$ rather than the usual homogeneous ISM. Several authors (Castor et al. 1975; Weaver 1977; Ramirez-Riuz et al. 2001) further found that beyond some typical radius of the wind bubble, the swept-up mass is too large to be pushed by the wind. As a result, the wind materials pile up at the edge of the wind bubble to form a density jump.  Lazzati et al. (2002) proposed that a density jump can lead to a rebrightening in the afterglow. In usual case, when the observing frequence is between the peak frequence ($\\nu_{\\rm m}$) and the cooling frequence ($\\nu_{\\rm c}$), the amplitude of the rebrightening should be proportional to the square root of the density contrast. However, if the density contrast is too high, the rebrightening will be weakened, since $\\nu_{\\rm c}$ will decrease and become less than the observing frequence. In the more detailed numerical simulations by Huang et al. (2007), no obvious rebrightenings can be discriminated when the density jump amplitude is set to 100, which suggests that the simple density jump model is not an ideal mechanism to produce the rebrightenings. In the recent studies by Nakar \\& Granot (2007) and van Eerten et al. (2009), a full treatment of the transient features at the jump moment, even including the thickness of the blastwave, was considered. Thus these studies should be a more accurate approximation to the reality. Interestingly, no obvious rebrightenings associated with the density jumps are found.  In short, a simple density jump model is difficult to explain the observed rebrightenings in GRB afterglows. On the other hand, the early afterglows of most GRBs exhibit flattish decays with $\\alpha < 0.8$, where $\\alpha$ is defined as $F_{\\nu} \\propto t^{-\\alpha}$. This is difficult to explain using the standard fireball model when the density profile is $\\rho \\propto r^{-2}$. Some previous works have suggested that the microphysics parameters may vary during the evolution of the fireball (Rossi \\& Rees 2003; Ioka et al. 2006; Fan \\& Piran 2006; Panaitescu et al. 2006; Granot, K\\\"{o}nigl \\& Piran 2006). We believe that this kind of variation could resolve these problems.  In this paper, we show that the observed rebrightenings in GRB afterglows can be well reproduced by assuming varying microphysics shock parameters associated with the wind bubble environments. The outline of our paper is as follows: in \\S 2 we introduce our model detailedly. We then numerically investigate the effects of various parameters on the afterglows in \\S 3, and reproduce the R-band and X-ray afterglow light curves of GRB 060206, GRB 070311 and GRB 071010A in \\S 4. Our discussion and conclusions are presented in \\S 5. We use an assumptive cosmology of $H_{\\rm 0} = 65$ ${\\rm km}$ ${\\rm s}^{-1}$ ${\\rm Mpc}^{-1}$, $\\Omega_{\\rm M} = 0.30$ and $\\Omega_{\\rm \\Lambda} = 0.70$ throughout the paper. ",
        "conclusions": "In the standard fireball model, it is assumed that the circum-burst medium density is constant or following a single $\\rho \\propto r^{-2}$ law. The microphysics shock parameters are also usually assumed invariable during the evolution of the fireball. These models can basically explaining many pre-$Swift$ GRB afterglows, whose spectrum and light curves can be approximated as broken power-law functions (Panaitescu \\& Kumar 2001a, 2001b; Yost et al. 2003). The launch of \\emph{Swift} satellite (Gehrels 2004) makes it possible to observe early afterglows of many GRBs in the first few hours after the trigger. Many remarkable and unexpected features are found, such as marked rebrightenings and the flattish decay phase in early GRB afterglows.  In fact, the standard fireball model is obviously too simplified. If long GRBs indeed originate from the death of massive stars, the environment of GRBs should have complex structures rather than have a constant or single power law profile. Some analytical (Castor et al. 1975; Weaver 1977; Ramirez-Riuz et al. 2005; Pe'er and Wijers 2006) and numerical (Ramirez-Riuz et al. 2001) studies suggest that the circum-burst density profile should be a wind bubble, associated with a few density jumps. On the other hand, the microphysics parameters, such as $\\xi_{\\rm e}$ and $\\xi_{\\rm B}$, may vary during the evolution of the fireball (Rossi \\& Rees 2003; Ioka et al. 2006; Fan \\& Piran 2006; Panaitescu et al. 2006; Granot, K\\\"{o}nigl \\& Piran 2006). Actually, the fast decrease of the cooling frequency $\\nu_c$ in some GRBs suggests that $\\xi_{\\rm B}$ may be evolving (Panaitescu et al. 2006). Some previous studies also suggest that $\\xi_{\\rm e}$ and $\\xi_{\\rm B}$ may be different for the forward shock and the reverse shock (Fan et al. 2002; Wei, Yan \\& Fan 2006), and may also be different for Region (1) and Region (2) (Gebdre et al. 2007; Kamble, Resmi \\& Misra 2007). So, $\\xi_{\\rm e}$ and $\\xi_{\\rm B}$ may depend on the strength of the shock and the environment.  In our study, we combine the wind bubble environments and the change of microphysics shock parameters together. Comparing with the standard fireball model, our model has three more parameters (i.e. $r_{\\rm wind}$, $\\alpha1$ and $\\alpha2$).  We find that this model can produce the observed rebrightenings and flattish decay in GRB afterglows successfully. We have shown that the observed rebrightenings in both the optical and X-ray afterglow light curves of GRB 060206, GRB 070311 and GRB 071010A can be well explained by our model.  We have also investigated the effects of various parameters on the light curves numerically. We can imagine that if the stellar wind produced by the progenitor is very strong, or the launching speed of the stellar wind $v_{\\rm w}$ is small, or the value of $r_{\\rm wind}$ is large enough, there will be no rebrightening within the usual observation time. So, basically the wind bubble environment can also give birth to a steadily decaying afterglow that shows no rebrightening.  In our work, we neglect the effect of the reverse shock. According to the study by Dai \\& Lu (2002), when an ultrarelativistic blast wave interacts with a density jump medium, the resulting reverse shock is relativistic only if the amplitude of the density jump is much larger than 21. In our model, the amplitude of the density jump is only 4, so the corresponding reverse shock is Newtonian. The emission from the Newtonian reverse shock is very weak, and can be omitted. Moreover, although we use an abrupt density jump, the actual increase of density may be gradual. In this situation, the emission from the reverse shock will even be much weaker.  Currently, a complete understanding of the microphysical processes in the relativistic shocks is still lacking. Using the derived microphysical parameters from GRB modeling, people may be able to get some constraints on the shock physics. The derived parameter values of $\\dot M$ and $n_{\\rm ISM}$ are also very important. They are closely related to the evolution and the environment of the massive star. They may give some hints on the characters of the progenitor, such as the initial main-sequence mass and the metallicity. So, from the modeling parameters, we can also obtain some useful information about GRB origin."
    },
    "0910/0910.4578_arXiv.txt": {
        "abstract": "The Cosmic Neutrino Background ({\\cnub}) anisotropy is calculated for massive neutrino states by solving the full Boltzmann equation. The effect of weak gravitational lensing, including the Limber approximation, is also derived for massive particles, and subsequently applied to the case of massive neutrinos. ",
        "introduction": "The Cosmic photon Microwave Background (CMB) is currently our main source of information about the physical content of the Universe. Observations of the CMB anisotropy provides detailed information about the curvature of the Universe, the matter content, and a plethora of other parameters \\cite{Komatsu:2008hk}.  Standard model physics likewise predicts the presence of a Cosmic Neutrino Background (\\cnub) with a well defined temperature of $T_\\nu \\sim (4/11)^{1/3} T_\\gamma$. While it remains undetected in direct experiments, the presence of the {\\cnub} is strongly hinted at in CMB data. The homogeneous C$\\nu$B component has been detected at the 4-5$\\sigma$ level in the WMAP data (see e.g.\\ \\cite{Komatsu:2008hk,Hamann:2007pi,de Bernardis:2007bu,Ichikawa:2008pz,Hamann:2008we,Popa:2008tb}). Furthermore, this component is known to be free-streaming, i.e. to have an anisotropic stress component consistent with what is expected from standard model neutrinos (see \\cite{Bashinsky:2003tk,Trotta:2004ty,Bell:2005dr,DeBernardis:2008ys,Basboll:2008fx,Hannestad:2004qu,Friedland:2007vv}). Finally the standard model neutrino decoupling history is also confirmed by Big Bang Nucleosynthesis (BBN), the outcome of which depends on both the energy density and flavour composition of the {\\cnub}.  While this indirect evidence for the presence of a {\\cnub} is important, a direct detection remains an intriguing, but almost impossible goal. The most credible proposed method is to look for a peak in beta decay spectra related to neutrino absorption from the {\\cnub} \\cite{Weinberg:1962zz,Cocco:2007za,Blennow:2008fh}, although many other possibilities have been discussed \\cite{Weiler:1982qy,Stodolsky:1974aq,Gelmini:2004hg,Ringwald:2004np,Fodor:2002hy,Duda:2001hd,Langacker:1982ih,Cabibbo:1982bb}. The neutrino absorption method was first investigated by Weinberg \\cite{Weinberg:1962zz}, based on the possibility that the primordial neutrino density could be orders of magnitude higher than normally assumed due to the presence of a large chemical potential. Although a large chemical potential has been ruled out because it is in conflict with BBN and CMB \\cite{Pastor:2001iu,Pastor:2008ti,Simha:2008mt,Wong:2002fa,Abazajian:2002qx}, the method may still work and recently there has been renewed interest in detecting the {\\cnub} using beta unstable nuclei.  Although the direct detection of the {\\cnub} is already very challenging, one might speculate on the possibility that in the more distant future anisotropies in the {\\cnub} will be detectable. For massless neutrinos the calculation of the {\\cnub} proceeds in a way which is almost identical to the standard CMB calculation. The massless {\\cnub} anisotropy spectrum was first presented in \\cite{Hu:1995fqa} and subsequently calculated in \\cite{Michney:2006mk} in a highly simplified way which contains some, but not all of the essential physics.  For massive neutrinos the calculation is much more complicated, and the {\\cnub} anisotropy is changed considerably: If the mass is sufficiently high, neutrino velocities can be as low as the escape velocities of galaxies. In this case the {\\cnub} is entirely determined by non-linear gravitational clustering. The current thermal velocity of a non-relativistic homogeneous neutrino background is given roughly by $\\langle v \\rangle \\sim 1500 \\, {\\rm km} \\, {\\rm s}^{-1} \\left(\\frac{0.1 \\, {\\rm eV}}{m_\\nu}\\right)$, which should be compared to gravitational streaming velocities which are up to $\\sim 1000 \\, {\\rm km} \\, {\\rm s}^{-1}$. For small masses (i.e.\\ non-degenerate, $m_\\nu \\lwig 0.1$ eV) it is possible to make a calculation which is analogous to what is done for the CMB. As will be explained later the {\\cnub} spectrum is shifted to larger angular scales, mainly because the much shorter distance to the neutrino last scattering surface changes the relation between angular scales and length scales. Furthermore the amplitude of the anisotropy greatly increases at small multipoles because the gravitational source term becomes much more important as neutrinos become increasingly non-relativistic at late times.  In addition to this change in the primary {\\cnub} spectrum, the effect of gravitational lensing is also very different from the case of massless neutrinos. As is the case for the primary spectrum, gravitational lensing becomes increasingly important at low $l$ as the mass increases. A detailed calculation of the lensing of massive neutrinos is presented in section 3. First, however, we derive the necessary equations for the primary {\\cnub} anisotropy in section 2 and present a numerical calculation of the {\\cnub} power spectra for different masses. In section 4 we combine the results of sections 2 and 3 to derive the form of the lensed massive {\\cnub}. Finally, section 5 contains a discussion of our results and our conclusions. ",
        "conclusions": "We have calculated the anisotropy of the {\\cnub} in linear theory which applies to neutrino masses of less than $\\sim 0.1$ eV. For massless neutrinos the power spectrum of {\\cnub} fluctuations closely resembles the usual CMB spectrum, but with the baryon-photon acoustic oscillations absent.  At high $l$ the neutrino spectra are almost identical, independent of the neutrino mass. The reason is that at high $l$ all neutrinos are dominated by free-streaming which in $l$-space has approximately the same impact for all masses.  For smaller $l$-values the anisotropy increases dramatically as the mass increases, because the gravitational source term becomes much more important at late times for massive particles. This initially increases the lowest multipoles but via the Boltzmann hierarchy the effect quickly propagates to higher $l$.  We then proceeded to calculate the effect of weak gravitational lensing for massive neutrinos and found it to be much stronger at low $l$, and correspondingly weaker at high $l$, as compared to the massless case. Finally we calculated the effect of lensing on the primary {\\cnub} spectra and found the effect to be unimportant (with relative changes at the per mille level up to $l \\sim 50$), but with some differences depending on the neutrino mass.  It is worth mentioning that any direct experimental measurement of the {\\cnub} anisotropy will most likely measure flavour states, not mass states. The actual anisotropy measured will therefore be a superposition of anisotropies for three different mass states, weighed with their individual flavour content.  We should finally again stress that our results are only valid for masses of $\\lwig 0.1$ eV. For higher masses linear perturbation theory breaks down because neutrino streaming velocities become comparable to the typical gravitational flow velocities so that a significant fraction of neutrinos are bound in structures. In this case the {\\cnub} spectrum must be found from $N$-body simulations of neutrino clustering \\cite{Brandbyge1}. This can also be seen from the fact that the anisotropy at low $l$ is a factor $\\sim 10^9$ higher for $m_\\nu=0.1$ eV than for massless neutrinos. Since the anisotropy for massless particles corresponds to $\\delta \\rho/\\rho \\sim 10^{-5}$ the corresponding $\\delta \\rho/\\rho$ for 0.1 eV neutrinos is of order one, indicating that perturbation theory breaks down."
    },
    "0910/0910.1731_arXiv.txt": {
        "abstract": "The mass accretion rate of transonic spherical accretion flow onto compact objects such as black holes is known as the Bondi accretion rate, which is determined only by the density and the temperature of gas at the outer boundary. A rotating accretion flow has angular momentum, which modifies the flow profile from the spherical Bondi flow, and hence its mass accretion rate, but most work on disc accretion has taken the mass flux to be a given with the relation between that parameter and external conditions left uncertain. Within the framework of a slim $\\alpha$ disk, we have constructed global solutions of the rotating, viscous hot accretion flow in the Paczy\\'{n}ski-Wiita potential and determined its mass accretion rate as a function of density, temperature, and angular momentum of gas at the outer boundary. We find that the low angular momentum flow resembles the spherical Bondi flow and its mass accretion rate approaches the Bondi accretion rate for the same density and temperature at the outer boundary. The high angular momentum flow on the other hand is the conventional hot accretion disk with advection, but its mass accretion rate can be significantly smaller than the Bondi accretion rate with the same boundary conditions. We also find that solutions exist only within a limited range of dimensionless mass accretion rate $\\dot{m} \\equiv \\dot{M}/\\dot{M}_B$, where $\\dot{M}$ is the mass accretion rate and $\\dot{M}_B$ the Bondi accretion rate: When the temperature at the outer boundary is equal to the virial temperature, solutions exist only for $0.05 \\la \\dot{m} \\le 1$ when $\\alpha=0.01$. We also find that the dimensionless mass accretion rate is roughly independent of the radius of the outer boundary but inversely proportional to the angular momentum at the outer boundary and proportional to the viscosity parameter, $\\dot{m} \\simeq 9.0\\ \\alpha \\lambda^{-1}$ when $0.1 \\la \\dot{m} \\la 1$, where the dimensionless angular momentum measure $\\lambda \\equiv l_{out}/l_B$ is the specific angular momentum of gas at the outer boundary $l_{out}$ in units of $l_B \\equiv GM/c_{s,out}$, $M$ the mass of the central black hole, and $c_{s,out}$ the isothermal sound speed at the outer boundary. ",
        "introduction": "The amount of mass gravitationally accreted to the compact objects such as black holes is determined by the conditions of gas around the compact objects. The case when gas has a spherically symmetric distribution, a polytropic pressure-density relation, and no angular momentum was solved by Bondi (1952), who found that the mass accretion rate for the transonic accretion is determined only by the density and the temperature of surrounding gas, both assumed to approach constant values far from the accreting objects. This rate, known as the Bondi accretion rate, is widely used as `the mass accretion rate' in a variety of accretion problems. As a function of the density $\\rho_\\infty$ and the isothermal sound speed $c_{s,\\infty}$ of gas at infinity, the Bondi rate for pure hydrogen gas is \\begin{equation}\\label{eq:Bondi_rate} \\dot{M}_B =  4 \\pi \\Lambda \\frac{(GM)^2 \\rho_\\infty}{\\gamma^{3/2}c_{s,\\infty}^3}, \\end{equation} where $G$ is the gravitational constant, $M$ the mass of the central object, and $\\gamma$ the adiabatic index of accreting gas. The constant $\\Lambda(\\gamma)$ is 0.25 for $\\gamma=5/3$ and 1.12 for $\\gamma=1$. If we define the Bondi radius as $r_B \\equiv GM/c_{s,\\infty}^2$, then $\\dot{M}_B = \\Lambda \\gamma^{-3/2} 4\\pi r_B^2 \\rho_\\infty c_{s,\\infty}$. Gravity starts to dominate over gas pressure inside the Bondi radius and the Bondi accretion rate is roughly the mass flux of gas infalling into a sphere of radius $r_B$ with a velocity equal to the sound speed at infinity.  However, many astrophysical accretion flows are expected to have certain amount of angular momentum, and the angular momentum will surely affect the flow properties, including the mass accretion rate. Surprisingly, the rate of mass accretion for these rotating, viscous accretion flows for given external boundary conditions, has not been studied yet. We might expect that if the rotational support is negligible at the Bondi radius, the effect of angular momentum to mass accretion rate would be small. But the accretion rate would diminish as the dimensionless measure of rotation $\\lambda \\equiv l_{out}/l_B$ approaches unity, where $l_{out}$ is the specific angular momentum of gas at the outer boundary and $l_B \\equiv r_B c_{s,\\infty} $ is the representative angular momentum expected at the Bondi radius.  Although the accretion problem can be reduced to local one in the limit of thin accretion disc where the radial velocity is negligible (Shakura \\& Sunyaev 1973), accretion flow is fundamentally global in the sense that the flow structure is only globally determined. As pointed out by Yuan (1999), the flow property can be qualitatively affected by the outer boundary conditions. The dependence of the mass accretion rate on boundary conditions can only be addressed by constructing global solutions. There have been many studies on the global solutions of rotating accretion flow (e.g., Narayan, Kato, \\& Honma 1997; Chen, Abramowicz, \\& Lasota 1997; Nakamura et al. 1997; Lu, Gu, \\& Yuan 1999) and the effects of outer boundary conditions on accretion flow (Yuan 1999; Yuan et al. 2000), but most have focused on the flow structures and emission properties for a given mass accretion rate, and the determination of how that rate is fixed by external circumstances remains to be addressed.  In this work, we focus on the mass accretion rate of the accretion flow, especially its dependence on the density, temperature, and the angular momentum of gas at the outer boundary. We do this by constructing the simplest global, transonic solutions within the slim disk formalism (Abramowicz et al. 1988). Slim disk formalism uses vertical integration to take into account the finite thickness of the accretion disk, and more importantly allows for the radial motion of the accretion flow and the critical point unlike the thin disk approximation (Shakura \\& Sunyaev 1973). ",
        "conclusions": "\\subsection{Flow properties}  Figure 1 shows typical solutions with different amount of angular momentum: high angular momentum (solid lines), intermediate angular momentum (dotted lines), and low angular momentum (dashed lines). By high angular momentum, we mean $\\lambda \\sim 1$, by low angular momentum, $\\lambda \\ll 1$, and by the intermediate angular momentum, between the two limits. The density and the temperature of gas at the outer boundary are the same for the three solutions: $T_{out} = 2.75\\times 10^{9} \\K$ and $\\rho_{out} = 1.5\\times 10^{-6} \\rho_0$ at $r_{out} = 10^3 r_{Sch}$. The dimensionless angular momenta of gas at the boundary and the mass accretion rate are: $\\lambda = 1.7$ and $\\dot{m} = 0.05$ for high angular momentum flow; $\\lambda = 0.27$ and $\\dot{m} = 0.386$ for intermediate angular momentum flow; $\\lambda = 0.14$ and $\\dot{m} = 0.643$ for low angular momentum flow. Most of the thermal energy is advected with the flow rather than radiated away, and temperature stays close to the virial value.  The flow with high angular momentum becomes supersonic at small radius $r_{cr} = 2.2 r_{Sch}$ (solid line in Fig. 1a) and the angular momentum (solie line in Fig. 1c) is close to or a few times lower than the Keplerian value (short-long dashed line in Fig. 1c). The radial velocity is significant but still smaller than the near free-fall velocity of spherical Bondi flow. This is a typical hot, radiatively inefficient advection-dominated accretion flow (ADAF) that has been extensively studied (e.g., Narayan \\& Yi 1994, 1995; Abramowicz et al. 1995; Narayan, Kato, \\& Honma 1997; Chen, Abramowicz, \\& Lasota 1997).  In spherical accretion, the flow passes the critical point at a much larger radius $r_{cr} \\sim r_B$, and the flow becomes supersonic inside $\\sim r_B$. But in rotating viscous accretion flow, a large part of the flow is subsonic well inside $r_B$ because the rotation of the flow balances out gravity, and the value of $r_{cr}$ depends on the angular momentum of the flow. A low angular momentum flow (dashed line in Fig. 1a) has a critical point at a much larger radius $r_{cr} = 17 r_{Sch}$ compared to the high angular momentum one. It has a larger radial infall velocity at the outer boundary as well because of smaller centrifugal force (Fig. 1b), and the mass accretion rate is higher than that of the high angular momentum flow. These characteristics show that this low angular momentum flow has more resemblance to the spherical flow than to the disk flow as is expected. This type of flow solution is first discovered by Yuan (1999) and its dynamical and thermal properties have been comprehensibly studied by Lu, Gu, \\& Yuan (1999) and Yuan et al. (2000).  The flow with intermediate angular momentum (dotted lines) has a critical point at $r_{cr} = 3.2 r_{Sch}$, and shows intermediate characteristics between the high and low angular momentum flow. The flow properties change from disklike to quasi-spherical as the angular momentum of the flow decreases. This is similar to the case of inviscid rotating accretion flow, which becomes disklike or quasi-spherical, depending on the angular momentum (Abramowicz \\& Zurek 1981). This transition of accretion flow from disklike to quasi-spherical in terms of critical radius position is described in detail by Yuan et al. (2000).  \\subsection{Mass accretion rate}  In spherical accretion, the outer boundary is naturally selected to be outside the critical point, where the density and the temperature of gas stay roughly constant. In rotating viscous flow, the outer boundary is not naturally associated with the critical point which is at much smaller radius. So we consider two choices for the outer boundary radius: one at the Bondi radius, $r_B(T_{out})$, for given $T_{out}$ and the other at a fixed radius regardless of $T_{out}$. The former choice is equivalent to setting the outer boundary at a virial temperature if we define the virial temperature as $T_{vir}(r) \\equiv GMm_p / (2 k r)$.  While original Bondi solutions are expressed in terms of the density and temperature of gas at infinity, our solutions start at a finite radius. But since the density and temperature of gas vary little from infinity to the Bondi radius in the Bondi solutions, we compare our solutions for given density and temperature of gas at finite $r_{out}$ against the Bondi solutions with the same density and temperature at infinity.  We first discuss the case where the outer boundary is set to be the Bondi radius, $r_{out} = GM/c_{s,out}^2 \\propto T_{out}^{-1}$ for given $T_{out}$, so that the solutions can be meaningfully compared with the Bondi solutions. The density at the outer boundary is either $\\rho_{out} = 1.5 \\times 10^{-6} \\rho_0$ or $1.5 \\times 10^{-9} \\rho_0$. Since the whole set of equations can be rescaled in density except the cooling function $q^-$, the mass accretion rate then is simply proportional to $\\rho_{out}$ when cooling is not important. In such case, the flow profile depends only on $\\dot m$ and not on the specific value of $\\rho_{out}$. However, as the mass accretion rate $\\dot M$ approaches 0.1 times the Eddington mass accretion rate, $\\dot{M}_{Edd} \\equiv L_{Edd}/c^2$, cooling becomes important as in spherical accretion (Park 1990). In such high $\\dot M/\\dot{M}_{Edd}$ case, the existence and the properties of the flow depend on the specific value of $\\rho_{out}/\\rho_0$.  Figure 2 shows the mass accretion rate of constructed solutions as a function of the angular momentum at the outer boundary, in $\\dot{m}$ versus $\\lambda$ plane. Symbols represent the solutions for different $T_{out}$: circles for $T_{out}=1.1\\times10^{10} \\K$ ($r_{out}=2.5\\times 10^2 r_{Sch}$), triangles for $T_{out}=5.5\\times10^9 \\K$ ($r_{out}=5.0\\times 10^2 r_{Sch}$), squares for $T_{out}=2.8\\times10^9 \\K$ ($r_{out}=1.0\\times 10^3 r_{Sch}$), pentagons for $T_{out}=1.1\\times10^9 \\K$ ($r_{out}=2.5\\times 10^3 r_{Sch}$), and hexagons for $T_{out}=5.5\\times10^8 \\K$ ($r_{out}=5.0\\times 10^3 r_{Sch}$). All these solutions are calculated for $\\rho_{out} = 1.5 \\times 10^{-6} \\rho_0$. The solutions for $T_{out}=1.1\\times10^7 \\K$ ($r_{out}=2.5\\times 10^5 r_{Sch}$), temperature suitable for the ISM in a galactic nucleus, have much larger $r_{out}$, and the bifurcation between subsonic to unphysical branch becomes even sharper. Since the critical mass accretion rate above which the hot, optically thin accretion disk does not exist also decreases as $r_{out}$ increases (Abramowicz et al. 1995) and the Bondi rate increases as $T_{out}$ decreases, we choose $\\rho_{out} = 1.5 \\times 10^{-9} \\rho_0$ so that $\\dot{M} /\\dot{M}_{Edd}$ is in a range where solutions can be found. The mass accretion rates of these solutions are shown as crosses in the center of Figure 2.  We find that regardless of $T_{out}$, the mass accretion rate $\\dot{m}$ decreases as the angular momentum $\\lambda$ increases. It approaches the Bondi accretion rate as $\\lambda$ decreases. This change of the mass accretion rate on the angular momentum is expected because given density and temperature at $r_{out}$, $\\rho$ and $H$ are fixed in the continuity equation (Eq. [\\ref{eq:continuity}]), and the mass accretion rate is determined solely by the radial velocity $v_r$. From equation (\\ref{eq:radial}), the radial velocity is determined by the difference between the gravity and the centrifugal accleration which is proportional to $l^2/r^3$ (at $r \\gg r_{Sch}$). Larger angular momentum causes smaller infall velocity, and the mass accretion rate decreases, and vice versa.  The mass accretion rate for the lowest angular momentum flow is quite close to the Bondi rate ($\\dot{m} \\sim 1$) while that for the highest angular momentum flow is roughly 20 times smaller than the corresponding Bondi rate ($\\dot{m} \\sim 0.05$). We also find that $\\dot{m}$ as a function of $\\lambda$ is rather insensitive to $T_{out}$, or equally $r_{out}$ for $\\dot{m} \\ga 0.1$: The slope of $\\dot m(\\lambda)$ has a tendency to become slightly flatter (from circles to crosses) as $T_{out}$ decreases but the difference is not large.  Since we expect the flow profile, especially the radial velocity, to depend on the specific value of $\\alpha$, we also calculate the solutions for different values of $\\alpha$. The mass accretion rates of $\\alpha=0.003$ are shown as a series of stars in the left side of Figure 2 and that of $\\alpha=0.03$ are shown in the right side of of Figure 2. Compared to the solutions for $\\alpha=0.01$ in the center denoted by crosses, lower $\\alpha$ flows have smaller mass accretion rate and higher $\\alpha$ flows larger mass accretion rate. Larger $\\alpha$ means stronger viscosity for given density and temperature, which causes larger radial velocity, and hence larger mass accretion rate, and vice versa for smaller $\\alpha$. The relation between $\\dot m$ and $\\lambda$  for $0.1 \\la \\dot{m} \\la 1$ can be approximated by \\begin{equation}\\label{eq:mdot_mB} \\dot{m}(\\lambda) = 0.09  \\left( \\frac{\\alpha}{0.01} \\right) \\lambda^{-1}. \\end{equation} The three straight lines (from left to right) in Figure 2 show this fitting function for $\\alpha=0.003$, 0.01, and 0.03, respectively.  The mass accretion rate of the self-similar ADAF that has density $\\rho_{out}$ at $r_{out}$ has the same dependency on $\\rho_{out}$ and $r_{out}$ as the Bondi rate $\\dot{M}_{B}$ (Narayan \\& Yi 1994). Hence, the dimensionless mass accretion rate of self-similar ADAF is simply $\\dot{m}_{ADAF} = 3.3 \\alpha$. Since ADAF generally refers to $\\lambda \\sim 1$ flow and $\\dot{m}(\\lambda=1) \\simeq 3 \\dot{m}_{ADAF}$, we find that our solutions have approximately three times higher mass accretion rates than those of self-similar ADAFs. The difference is probably due to the fact that the density and velocity of global ADAFs do not exactly follow the self-similar form, especially near the outer boundary (Narayan, Kato, \\& Honma 1997; Chen, Abramowicz, \\& Lasota 1997).  For given $T_{out}$, the lower limit on the mass accretion rate is determined by the Keplerian angular momentum barrier. A smaller mass accretion rate demands smaller radial infall velocity. From equation (\\ref{eq:angular}), smaller $|v_r|$ requires larger $\\Omega$, and eventually $\\Omega$ becomes larger than the Keplerian value $\\Omega_K$ at $r_{out}$ for too small mass accretion rate. Gas cannot accrete in steady-state when the angular momentum is much larger than the Keplerian one, and therefore no solution exists when $\\lambda \\ga 3$. Although we could determine $\\dot m$ below $\\dot{m} \\simeq 0.1$ from shooting until $\\lambda$ approaches $\\sim 1.5 - 3$ or 3, in many cases we failed to construct the full transonic solutions down to the inner boundary. The difficulty is much more severe for small $\\alpha$: the bifurcation between subsonic and unphysical occurs at a large radius and computing accuracy is not enough to integrate all the way down to the horizon. So we have limited confidence in the solutions below $\\dot{m} \\la 0.2$ for $\\alpha = 0.003$ and 0.03. The upper limit on the mass accretion rate on the other hand is determined by the Bondi accretion rate itself. The lowest angular momentum flow has its mass accretion rate already quite close to the Bondi rate, which is the mass accretion rate for zero angular momentum flow. Thus any higher mass accretion flow will end up on the unphysical branch of solutions in spherical Bondi flow (type III), with minimum effect due to the angular momentum.  We can think of extending the velocity versus radius diagram of Bondi (1952) to a similar one with the angular momentum effect added. What angular momentum does is to lower the critical mass accretion rate for the transonic flow. As the angular momentum of the flow increases, the critical mass accretion rate decreases, and the rotating accretion flow accretes transonically at a critical mass accretion rate lower than the original Bondi accretion rate. On the other hand, if the angular momentum of the flow decreases, the critical mass accretion rate increases up to the Bondi accretion rate which is the critical mass accretion rate for zero angular momentum flow, above which no transonic steady-state rotating accretion flow exists. Since equation (\\ref{eq:mdot_mB}) gives $\\dot{m} > 1$ for $\\lambda < 9 \\alpha$ and $\\dot{m}$ cannot be greater than unity, it looks like there is no solution for $\\lambda < 9 \\alpha$. However, we are using slim disk formulation with a simple viscosity prescription to describe what is really a two-dimensional quasi-spherical flow, especially for lowest angular momentum flow; so the non-existence of solutions below the lower limit on $\\lambda$ is probably caused by the approximate physical descriptions of the flow, especially the overestimated viscosity for the low or no angular momentum flow. Surely, we will have a spherical Bondi flow when the angular momentum is zero. Therefore, we expect that steady flow probably exists even below $\\lambda \\sim 9 \\alpha$ and \\begin{equation}\\label{eq:mdot_zerolambda} \\dot{m} \\simeq 1 \\quad {\\rm for} \\quad \\lambda \\la 9 \\alpha . \\end{equation} We conclude that the steady-state, transonic hot accretion flow exists only when the mass accretion rate is within a certain range or the angular momentum at the outer boundary is below a certain maximum value. For our specific choice of parameters and conditions, the range in the mass accretion rate is \\begin{equation}\\label{eq:mdot_range} \\dot{m}_{cr} \\la \\dot{m} \\leq  1, \\end{equation} where $\\dot{m}_{cr}=\\dot{m}(\\lambda_{cr})$ is the mass accretion rate of the flow with maximum possible angular momentum at the outer boundary. Our calculations suggest $\\lambda_{cr} \\la 3$ and $\\dot{m}_{cr} \\sim 0.05$ for $\\alpha = 0.01$.  Now we discuss flows with the outer boundary temperature different from the virial temperature. For example, when the cool, outer geometrically thin disk is attached to the hot inner accretion disk, the temperature at the outer boundary will be much lower than the virial temperature (Narayan, Kato, \\& Honma 1997). So we fix $r_{out} = 10^3 ~r_{Sch}$ but vary $T_{out}$. The outer boundary density is the same at $\\rho_{out} = 1.5 \\times 10^{-6} \\rho_0$. Figure 3 shows the mass accretion rate again in $\\dot{m}$ versus $\\lambda$ plane. Symbols are the same as in Figure 2, except circles for $T_{out}=3.6\\times10^9 \\K$, triangles for $T_{out}=2.2\\times10^9 \\K$, squares for $T_{out}=1.1\\times10^9 \\K$, pentagons for $T_{out}=5.5\\times10^8 \\K$, and hexagons for $T_{out}=2.0\\times10^8 \\K$.  Again, solutions exist only in a limited range of $\\dot{m}$, or equally $\\lambda$. But this time both upper and lower limits on $\\dot{m}$ vary with $T_{out}$. The reason for the difference is that $r_{out}$ is fixed and not related to $T_{out}$ whereas it is related to $T_{out}$ in Bondi solutions. As in previous choice of boundary, the lower limit on $\\dot{m}$ is from the Keplerian angular momentum barrier. However, the upper limit is determined for a different reason in this case. As the mass accretion rate for given boundary conditions increases, the radial infall velocity at the outer boundary should increase. But the increasing radial infall velocity eventually becomes larger than the sound speed at the boundary. The whole flow becomes supersonic from the outer boundary to the innermost radius. Since we are looking for transonic solutions that uniquely determine the mass accretion rate, we disregard these supersonic branch of solutions, and no transonic solutions exist above certain upper limit on $\\dot{m}$. This upper limit on $\\dot{m}$ decreases as $T_{out}$ decreases because the sound speed decreases with $T_{out}$ and the flow becomes supersonic at smaller radial velocity, or mass accretion rate, causing rapid decrease of the upper limit on $\\dot{m}$ with decreasing $T_{out}$.  The dependence of $\\dot{m}$ on $\\lambda$ is also different. Again, $\\dot{m}$ decreases as $\\lambda$ increases, but is not inversely proportional to  $\\lambda$, nor is independent of $T_{out}$. Cooler flow, i.e., lower $T_{out}$, has lower mass accretion rate $\\dot{m}$ for a given $\\lambda$. When $\\lambda = 0.1$, $\\dot{m}$ for $T_{out}=2.2\\times10^9 \\K$ is $\\sim 1$ while that for $T_{out}=2.0\\times10^8 \\K$ is as low as $\\sim 10^{-4}$. This means that if gas is injected into a given radius with near-Keplerian angular momentum but with temperature much below the virial temperature at that radius, the mass accretion rate can be a few orders of magnitude below the Bondi accretion rate."
    },
    "0910/0910.0117_arXiv.txt": {
        "abstract": "We use Monte-Carlo Markov chain techniques to constrain acceptable parameter regions for the Munich L-Galaxies semi-analytic galaxy formation model.  Feedback from active galactic nuclei (AGN) is required to limit star-formation in the most massive galaxies. However, we show that the introduction of tidal stripping of dwarf galaxies as they fall into and merge with their host systems can lead to a reduction in the required degree of AGN feedback.  In addition, the new model correctly reproduces both the metallicity of large galaxies and the fraction of intracluster light. ",
        "introduction": "This paper describes the implementation of a model for dwarf galaxy disruption within the semi-analytic framework of \\citet[hereafter DLB07]{LuB07}.  The original model was suggested by \\citet{HBT08} as a way of both reducing the excess of dwarf galaxies and creating the intracluster light (ICL); however this was implemented \\emph{a posteriori}, acting only to reduce the dwarf population at the current day.  The new model follows the stripping of dwarfs as they fall into the halos of their parent galaxy, thus gradually reducing their mass: affecting their infall rates, increasing the time-scale for, and decreasing the magnitude of, the merger with the central object.  For the purposes of these conference proceedings, the most important result is that the masses of the black holes are reduced, with a corresponding reduction in the level of feedback of AGN energy into the interstellar medium.  This should be read in conjunction with the paper by Chris Short in this volume that investigates the degree to which the accretion energy is needed to provide feedback into the intracluster medium (ICM) in order to provide the observed entropy excess in clusters. ",
        "conclusions": "The black hole growth model presented here provides a real challenge for models that require AGN heating to raise the entropy of the ICM. Even using the original DLB07 model, with its higher black hole masses, 35 per cent of the available rest mass energy is required (see article by Chris Short in this volume, also \\citet{ShT09}). Similar conclusions have been reached by \\citet{BMB08} using the Durham GALFORM semi-analytic model.  \\begin{theacknowledgments}  This work was undertaken using the Virgo Consortium cluster of computers, COSMA.  BH acknowledges the support of his PhD scholarship from the Portuguese Science and Technology Foundation which supported him for most of the time while this work was developed.  PAT was supported by an STFC rolling grant. \\end{theacknowledgments}"
    },
    "0910/0910.5218_arXiv.txt": {
        "abstract": "We calculate the cosmic microwave background (CMB) bispectrum due to inhomogeneous reionization. We calculate all the terms that can contribute to the bispectrum that are products of first order terms on all scales in conformal Newtonian gauge. We also correctly account for the de-correlation between the matter density and initial conditions using perturbation theory up to third order. We find that the bispectrum is of local type as expected. For a reasonable model of reionization, in which the universe is completely ionized by redshift $z_{ri}\\sim 8$ with optical depth to the last scattering surface $\\tau_0=0.087$ the signal-to-noise ratio for detection of the CMB temperature bispectrum is $\\rm{S/N}\\sim 0.1$ and confusion in the estimation of primordial non-Gaussianity is  $f_{NL}\\sim -0.1$. For an extreme model with $z_{ri}\\sim 12.5$ and $\\tau_0=0.14$ we get $\\rm{S/N}\\sim 0.5$ and $f_{NL}\\sim -0.2$. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.2865_arXiv.txt": {
        "abstract": "Densely-packed, all-digital aperture arrays form a key area of technology development required for the Square Kilometre Array (SKA) radio telescope.  The design of real-time signal processing systems for digital aperture arrays is currently a central challenge in pathfinder projects worldwide. We describe an hierarchical, frequency-domain beamforming architecture for synthesising a sky beam from the wideband antenna feeds of digital aperture arrays. ",
        "introduction": "Densely-packed, all-digital aperture arrays form a key area of technology development required for the Square Kilometre Array (SKA) radio telescope. The design of real-time signal processing systems for aperture arrays is certainly one of the most challenging tasks in the design of next-generation instruments. In this paper, we describe the design and implementation of an hierarchical beamforming architecture for wideband radio telescopes.\\\\  This document is organised as follows. Chapter I is an introduction to the problem domain. Chapter II describes the theory and process of `beamforming' for aperture arrays used in radio astronomy. Chapter III describes the design and implementation of a frequency-domain beamformer. Chapter IV  describes calibration procedures for the array, and is followed by results and conclusions.\\\\  \\subsection{Aperture Phased Arrays} An aperture phased array is a collection of fixed, non-directional antennas whose outputs are combined to synthesise an effective directional antenna. This synthesised antenna is pointed electronically; significantly the physical component antennas do not ever move. The electronic combination of signals with beamforming techniques performs the same spatial filtering task as a parabolic dish would in a traditional radio telescope dish. Figure \\ref{BeamformingIllustration} illustrates the basic principle of this operation.\\\\  The flexibility, survey speed, system-level autonomy and scalability provided by antenna arrays make them extremely attractive for next-generation radio instruments, for use, in our application, in radio astronomy, but also applicable to other wideband detection and communication instruments. Future radio instruments will certainly be dependent on the availability of high-performance signal processing systems such as these. Since aperture arrays are seen as a critical antenna technology to achieve the required sky survey speed in future radio telescopes (see, for example, \\cite{2006SKA.memo.81} and \\cite{2007skads.memo.FeA}), the development of a scalable signal processing system is of great importance.\\\\  \\begin{figure} \\begin{center} \\includegraphics[width=70mm]{BeamformingDiagNoBackg.pdf} \\caption {An illustration of the electronic processing of signals received from an array of antennas for concentrating the signal from a certain direction. This is the basic principle of the phased aperture array: different delay configurations allow beams to be formed in different directions of \\(\\theta\\)} \\label{BeamformingIllustration} \\end{center} \\end{figure} ",
        "conclusions": "In this paper we propose that digital aperture arrays are a fast, useful and elegant collector technology for GHz radio astronomy, that the development of high-performance signal processing systems is critical to the success of next-generation radio instruments, and have described the design of a 4-element frequency-domain aperture array beamformer.\\\\"
    },
    "0910/0910.4391_arXiv.txt": {
        "abstract": "If dark energy (DE) couples to neutrinos, then there may be apparent violations of Lorentz/{\\it CPT} invariance in neutrino oscillations.  The DE-induced Lorentz/{\\it CPT} violation takes a specific form that introduces neutrino oscillations that are energy independent, differ for particles and antiparticles, and can lead to novel effects for neutrinos propagating through matter.  We show that ultra-high-energy neutrinos may provide one avenue to seek this type of Lorentz/{\\it CPT} violation in $\\nu_\\mu$-$\\nu_\\tau$ oscillations, improving the current sensitivity to such effects by seven orders of magnitude. Lorentz/{\\it CPT} violation in electron-neutrino oscillations may be probed with the zenith-angle dependence for high-energy atmospheric neutrinos. The ``smoking gun,'' for DE-neutrino coupling would, however, be a dependence of neutrino oscillations on the direction of the neutrino momentum relative to our peculiar velocity with respect to the CMB rest frame.  While the amplitude of this directional dependence is expected to be small, it may nevertheless be worth seeking in current data and may be a target for future neutrino experiments. ",
        "introduction": "The accelerated cosmic expansion \\cite{SNIa} poses difficult questions for theoretical physics~\\cite{Copeland:2006wr,Caldwell:2009ix,Silvestri:2009hh}. Is it simply due to a cosmological constant?  Is some new negative-pressure dark energy (DE) required?  Is general relativity modified at large distance scales? The major thrust of the empirical assault on these questions has been to determine whether the expansion history and growth of large-scale structure are consistent with a cosmological constant or require something more exotic \\cite{DETF}.  However, it may be profitable to explore whether there are other experimental consequences of the new physics---which we collectively refer to as DE, although it may involve a modification of gravity rather than the introduction of some new substance---responsible for accelerated expansion. If cosmic acceleration is due to a cosmological constant (i.e., if general relativity is valid and the equation-of-state parameter is $w=-1$), then the vacuum is Lorentz invariant.  If, however, something else is going on, then the ``vacuum'' has a preferred frame: the rest frame of the cosmic microwave background (CMB).  If, moreover, dark energy couples somehow to standard-model particles, then there may be testable (apparent) violations of Lorentz invariance.  For example, if DE is coupled to the pseudoscalar $F\\tilde F$ of electromagnetism \\cite{Carroll:1998zi}, there may be a ``cosmological birefringence'' that rotates the linear polarization of cosmological photons; CMB searches for such a rotation \\cite{Lue:1998mq} constrain this rotation to be less than a few degrees~\\cite{Feng:2006dp}.  Here we explore DE-induced Lorentz/{\\it CPT}-violating effects in the neutrino sector.  We show that the form of a Lorentz-violating coupling between neutrinos and dark energy is highly restricted under fairly general assumptions.\\footnote{The coupling of neutrinos to dark energy has also been considered in the context of ``mass-varying neutrinos''~\\cite{MaVaN}, but that implementation of the DE-neutrino coupling does not lead to the type of Lorentz/{\\it CPT}-violating effects we discuss here.} The coupling engenders an additional source for neutrino mixing (e.g., Ref.~\\cite{Gu:2005eq}), resulting in neutrino oscillations with a different energy dependence than vacuum oscillations and different oscillation probabilities for neutrinos and antineutrinos.  While similar Lorentz/{\\it CPT}-violating oscillations have been considered before~\\cite{bargeretal,CPT_nu,Kostelecky2004}, we emphasize here that cosmic acceleration dictates a specific form for such effects.  Data from Super-Kamiokande and K2K~\\cite{Gonzalez-Garcia2004} and AMANDA/IceCube~\\cite{Abbasi2009} already tightly constrain {\\it CPT}-violating parameters for $\\nu_\\mu$-$\\nu_\\tau$ mixing, and those from solar-neutrino experiments and KamLAND~\\cite{Bahcall2002} do so for $\\nu_e$-$\\nu_\\mu$ mixing. However, the effects of DE-induced {\\it CPT} violation become more significant at higher energies \\cite{Hooper:2005}.  Here we show that next-generation measurements of ultra-high-energy neutrinos produced by spallation of ultra-high-energy cosmic rays will increase the sensitivity to {\\it CPT}-violating $\\nu_\\mu$-$\\nu_\\tau$ oscillations by seven orders of magnitude. We also show that these {\\it CPT}-violating couplings may lead to novel effects in the zenith-angle dependence for atmospheric neutrinos in the $\\sim100$~GeV range.  While such {\\it CPT}-violating effects, if detected, could be attributed simply to intrinsic {\\it CPT} violation in fundamental physics, not related to DE, a DE-neutrino coupling further predicts a directional effect: the neutrino-mixing parameters depend on the neutrino propagation direction relative to our peculiar velocity with respect to the CMB rest frame.  While this signature will likely remain elusive even to next-generation experiments, it would, if detected, be a ``smoking gun'' for DE beyond a cosmological constant.  It is therefore worth considering as a long-range target for future neutrino experiments.  It may also be worthwhile to search current data in case an implementation of DE-neutrino coupling different from that we consider here leads to a different energy dependence for these directional effects.  We therefore work out explicitly the directional dependence to aid experimentalists who may wish to look for such correlations in current data.  Below, we first derive in Sec.~\\ref{sec:formalism} the form of the Lorentz/{\\it CPT} violation allowed by a DE-neutrino coupling and discuss the resulting neutrino-oscillation physics. In Sec.~\\ref{sec:cosmogenic} we apply the formalism to cosmogenic ultra-high-energy neutrinos, and obtain projected sensitivities of future detectors to these effects in $\\nu_\\mu$-$\\nu_\\tau$ oscillations.  In Sec.~\\ref{resonance} we discuss matter-induced effects for $\\nu_e$ oscillations in high-energy atmospheric neutrinos in the presence of Lorentz-invariance--violating mixings. Concrete formulas for the directional dependence on oscillation probabilities are given in Sec.~\\ref{sec:direction}. Finally, we discuss some theoretical implications in Sec.~\\ref{implications} and summarize and conclude in Sec.~\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  We studied the implications of an interaction between dark energy and neutrinos for neutrino oscillations.  The most general Lorentz/{\\it CPT}-violating term induced by dark energy (DE) takes the form $(a_L)^\\mu \\bar\\nu \\gamma_\\mu (1-\\gamma_5) \\nu$, where $(a_L)^\\mu$ is a four-vector normal to the CMB rest frame.  This introduces a new source for neutrino oscillations that are energy independent and different for neutrinos and antineutrinos.  Furthermore, the motion of the Earth with respect to the cosmic rest frame induces a directional dependence in the oscillation probabilities.  The current best limits to the DE-neutrino coupling we considered are obtained from atmospheric- and accelerator-neutrino experiments for $\\nu_\\mu$-$\\nu_\\tau$ mixing, and from solar and reactor experiments for $\\nu_e$-$\\nu_\\mu$ mixing.  However, the higher the neutrino energy, the more prominent the effect of the DE-neutrino interaction.  We therefore considered in this paper cosmogenic ultra-high-energy (energies of $10^{17}$--$10^{19}$ eV) neutrinos produced by the interaction of ultra-high-energy cosmic rays with CMB photons. We showed that future experiments targeting these neutrinos will improve the sensitivity to a DE-neutrino interaction by seven orders of magnitude, down to $m_{\\rm eff} \\sim 10^{-30}$ GeV compared with the current upper bound $m_{\\rm eff} \\alt 5\\times 10^{-23}$ GeV (Fig.~\\ref{fig:GZK}).  This corresponds to a sensitivity to an energy scale as large as $\\sim$10$^6$ GeV for the DE-neutrino interaction.  We then showed that the interplay of DE- and matter-induced neutrino mixing could induce a novel zenith-angle dependence for $\\nu_e$ oscillations in atmospheric neutrinos.  This effect may extend the sensitivity to Lorentz/{\\it CPT}-violating parameters in the $\\nu_e$ by roughly three orders of magnitude.  The real smoking gun of a DE-neutrino interaction (as opposed to some other origin for Lorentz/{\\it CPT} violation) would be a directional dependence of the oscillation probabilities.  The notion that Lorentz violation may give rise to a directional dependence is not new (e.g., Ref.~\\cite{Kostelecky:2004hg}) and searches for directional dependence in neutrino experiments have already been carried out (e.g., Ref.~\\cite{Adamson:2008ij}), but prior work has considered Lorentz-violating parameters introduced in an {\\it ad hoc} manner and/or tested for direction-dependent effects in a Sun-centered inertial frame.  We emphasize here that cosmic acceleration suggests that we seek a specific form of Lorentz violation, that where the preferred frame is aligned with the CMB rest frame. Even though such a signal is expected to be small, it is still worth seeking in existing and future experimental data.  We have not discussed specific models for a DE-neutrino interaction, beyond an illustrative toy model, but it may be interesting to do so (see also Ref.~\\cite{stephon}).  The theoretical motivation to expect such a coupling may admittedly be slim.  However, we are at square one in our understanding of DE, and such a coupling is no less likely to be expected, perhaps, than any of the many other manifestations of new cosmic-acceleration physics that have been considered.  Discovery of Lorentz/{\\it CPT}-violating effects would be extremely important, even if not attributable directly to dark energy.  A directional dependence, if discovered, would be absolutely remarkable, as it would provide moreover clear evidence that there is more to cosmic acceleration than simply a cosmological constant."
    },
    "0910/0910.1327_arXiv.txt": {
        "abstract": "Spectral evolution models are a widely used tool for determining the stellar content of galaxies. I provide a review of the latest developments in stellar atmosphere and evolution models, with an emphasis on massive stars. In contrast to the situation for low- and intermediate-mass stars, the current main challenge for spectral synthesis models are the uncertainties and rapid revision of current stellar evolution models. Spectral libraries, in particular those drawn from theoretical model atmospheres for hot stars, are relatively mature and can complement empirical templates for larger parameter space coverage. I introduce a new ultraviolet spectral library based on theoretical radiation-hydrodynamic atmospheres for hot massive stars. Application of this library to star-forming galaxies at high redshift, i.e., Lyman-break galaxies, will provide new insights into the abundances, initial mass function and ages of stars in the very early universe. ",
        "introduction": "Spectral evolution models attempt to reproduce the observed spectra of systems ranging in sizes from star clusters to luminous galaxies by numerically combining models or observations of star formation, stellar evolution, and stellar spectra. \\cite[Tinsley (1968)]{Tinsley68} is generally credited for initiating the field, and \\cite[Charlot \\& Bruzual's (1991)]{ChaBru91} introduction of the isochrone synthesis technique validated this approach for ages up to a Hubble time. The success of synthetic spectral models echoes the increased availability of computers in astronomy. The power of this technique has been amplified over the past decade by the growing importance of the internet and databases for storage, query, and distribution of spectral templates and related models. This trend is revolutionizing the field in a way similar to Tinsley's and Charlot's \\& Bruzual's achievements.  In this talk I will concentrate less on the technical but more on some of the astrophysical aspects of spectral evolution models. After a brief discussion of the main ingredients in the model I will address the two major components: stellar evolution models and stellar libraries. The emphasis will be on {\\em massive stars}, for which stellar evolution still is the major challenge, whereas reliable stellar libraries are becoming available in larger and larger numbers. This situation is somewhat opposite to the case of low-mass stars. I will introduce a new ultraviolet (UV) spectral library constructed from a grid of radiation-hydrodynamic models and discuss its potential for interpreting the rest-frame UV spectra of Lyman-break galaxies. Detecting the first stellar generations and modeling and understanding their spectra will be a major effort for the next decade. ",
        "conclusions": ""
    },
    "0910/0910.4502_arXiv.txt": {
        "abstract": "We investigate the viscous two temperature accretion disc flows around rotating black holes. We describe the global solution of accretion flows with a sub-Keplerian angular momentum profile, by solving the underlying conservation equations including explicit cooling processes selfconsistently. Bremsstrahlung, synchrotron and inverse Comptonization of soft photons are considered as possible cooling mechanisms. We focus on the set of solutions for sub-Eddington, Eddington and super-Eddington mass accretion rates around Schwarzschild and Kerr black holes with a Kerr parameter $0.998$. It is found that the flow, during its infall from the Keplerian to sub-Keplerian transition region to the black hole event horizon, passes through various phases of advection -- general advective paradigm to radiatively inefficient phase and vice versa. Hence the flow governs much lower electron temperature $\\sim 10^8-10^{9.5}$K, in the range of accretion rate in Eddington units $0.01\\lsim\\mdot\\lsim 100$, compared to the hot protons of temperature $\\sim 10^{10.2}-10^{11.8}$K. Therefore, the solution may potentially explain the hard X-rays and $\\gamma$-rays emitted from AGNs and X-ray binaries. We then compare the solutions for two different regimes of viscosity and conclude that a weakly viscous flow is expected to be cooling dominated, particularly at the inner region of the disc, compared to its highly viscous counter part which is radiatively inefficient. With all the solutions in hand, we finally reproduce the observed luminosities of the under-fed AGNs and quasars (e.g. Sgr~$A^*$) to ultra-luminous X-ray sources (e.g. SS433), at different combinations of input parameters such as mass accretion rate, ratio of specific heats. The set of solutions also predicts appropriately the luminosity observed in the highly luminous AGNs and ultra-luminous quasars (e.g. PKS~0743-67). ",
        "introduction": "The cool Keplerian accretion disc (Pringle \\& Rees 1972; Shakura \\& Sunyaev 1973; Novikov \\& Thorne 1973) was found to be inappropriate to explain observed hard X-rays, e.g. from Cyg~X-1 (Lightman \\& Shapiro 1975). It was argued that secular instability of the cool disc swells the optically thick, radiation dominated region to a hot, optically thin, gas dominated region resulting in hard component of spectrum $\\sim 100$KeV (Thorne \\& Price 1975; Shapiro, Lightman \\& Eardley 1976). This region is strictly of two temperatures with electron and ion temperatures respectively $\\sim 10^9$K and $\\sim 5\\times 10^{11}$K which confirms that cool, one temperature, pure Keplerian accretion solution is not unique. Indeed Eardley \\& Lightman (1975) found that a Keplerian disc is unstable due to thermal and viscous effects when viscosity parameter $\\alpha$ (Shakura \\& Sunyaev 1973) is constant. Later Eggum et al. (1985) showed by numerical simulations that the Keplerian disc with a constant $\\alpha$ collapses.  Around eighties, therefore, the idea of two component accretion disc started floating around. For example, Paczy\\'nski \\& Wiita (1980) described a geometrically thick regime of the accretion disc in the optically thick limit, while Rees et al. (1982) introduced accretion torus in the optically thin limit. Moreover the idea of sub-Keplerian, transonic accretion was introduced by Muchotrzeb \\& Paczy\\'nski (1982), which was later improved by other authors (Chakrabarti 1989, 1996; Mukhopadhyay 2003). Other models were proposed by e.g. Gierli\\'nski et al. (1999), Coppi (1999), Zdziarski et al. (2001), including a secondary component in the accretion disc. On the other hand, Narayan \\& Yi (1995) introduced a two temperature disc model in the regime of inefficient cooling resulting in a vertical thickening of the hot disc gas. Here the pressure forces are expected to become important in modifying the disc dynamics which is likely to be sub-Keplerian. Other models with similar properties were proposed by, e.g., Begelman (1978), Liang \\& Thompson (1980), Rees et al. (1982), Eggum, Coroniti \\& Katz (1988). Abramowicz et al. (1988) proposed a height-integrated disc model, namely ``slim disc\", having high optical depth of the accreting gas at super-Eddington accretion rate such that the diffusion time is longer than the viscous time. The model was further applied to study the thermal and viscous instabilities in optically thick accretion discs (Wallinder 1991; Chen \\& Taam 1993).  Shapiro, Lightman \\& Eardley (1976) initiated a two temperature Keplerian accretion disc at a low mass accretion rate which is optically thin and significantly hotter than the single temperature Keplerian disc of Shakura \\& Sunyaev (1973). The optically thin hot gas cools down through the bremsstrahlung and inverse-Compton processes and could explain various states of Cyg~X-1 (Melia \\& Misra 1993). Similarly, the ``ion torus\" model by Rees et al. (1982) was applied to explain AGNs at a low mass accretion rate. However, the two temperature model solutions by Shapiro, Lightman \\& Eardley (1976) appear thermally unstable. Narayan \\& Popham (1993) and subsequently Narayan \\& Yi (1995) showed that introduction of advection may stabilize the system. However, the solutions of Narayan \\& Yi (1995), while of two temperatures, could explain only a particular class of hot systems with inefficient cooling mechanisms. They also described the hot flow based on the assumption of ``self similarity\" which is just a ``plausible choice\". They kept the electron heating decoupled from the disc hydrodynamical computations which merely is an assumption. Later on, the solutions were attempted to generalize by Nakamura et al. (1997), Manmoto et al. (1997), Medvedev \\& Narayan (2001), relaxing efficiency of cooling into the systems, but concentrating only on specific classes of solutions. On the other hand, Chakrabarti \\& Titarchuk (1995) and later Mandal \\& Chakrabarti (2005) modeled two temperature accretion flows around Schwarzschild black holes in the general ``advective paradigm\", emphasizing possible formation of shock and its consequences therein. However, they also did not include the effect of electron heating self-consistently into the hydrodynamical equation, and thus the hydrodynamical quantities do not get coupled to the rate of electron heating (see also Rajesh \\& Mukhopadhyay 2009).  In the present paper, we model a selfconsistent accretion flow in the regime of two temperature transonic sub-Keplerian disc (see also Sinha, Rajesh \\& Mukhopadhyay 2009; a brief version of the present work, but around Schwarzschild black holes). We consider all the hydrodynamical equations of the disc along with thermal components and solve the coupled set of equations selfconsistently. We neither restrict to the advection dominated regime nor the self-similar solutions. We allow the disc to cool selfconsistently according to the thermo-hydrodynamical evolution and compute the corresponding cooling efficiency factor as a function of radial coordinate. We investigate that when does the disc switch from the radiatively inefficient nature to general advective paradigm and vice versa.  In order to implement our model to explain observed sources, we focus on the under-luminous AGNs and quasars (e.g. Sgr~$A^{*}$), ultra-luminous quasars and highly luminous AGNs (e.g. PKS~0743-67) and ultra-luminous X-ray (ULX) sources (e.g. SS433), when the last items are likely to be the ``radiation trapped'' accretion discs. While the first two cases correspond to respectively sub-Eddington and super-Eddington accretion flows around supermassive black holes, the last case corresponds to super-Eddington accretors around stellar mass black holes.  In the next section, we discuss the model equations describing the system and the procedure to solve them. Subsequently, we discuss the two temperature accretion disc flows around stellar mass and supermassive black holes, respectively in \\S3 and \\S4, for both sub-Eddington, Eddington and super-Eddington accretion rates. Section 5 compares the disc flow of low Shakura-Sunyaev (1973) $\\alpha$ with that of high $\\alpha$ and then between the flows around co and counter rotating black holes. Then we discuss the implications of the results with a summary in \\S6. ",
        "conclusions": "We model the two temperature accretion flow, particularly around black holes, combining the equations of conservation and comprehensive cooling processes. We consider self-consistently the important cooling mechanisms: bremsstrahlung, synchrotron and inverse Comptonization due to synchrotron photons, where ions and electrons are allowed to have different temperatures. As matter falls in, hot electrons cool through the various cooling mechanisms, particularly by the synchrotron emission when the magnetic field is high. This is particularly the case for the flow around stellar mass black holes where the magnetic field may also act as a boost in transporting the angular momentum. However, in the present paper, we do not consider such processes in detail, rather stick with the standard $\\alpha$-prescription.  By solving a complete set of disc equations we show that in general the disc system exhibits GAAF. However, in certain circumstances GAAF becomes radiatively inefficient, depending on the flow parameters and hence efficiency of cooling mechanisms. Transitions from GAAF to radiatively inefficient flow and vice versa are clearly explained by the cooling efficiency factor $f$, shown in each model cases. While the previous authors, who proposed ADAF (Narayan \\& Yi 1994, 1995), especially restricted with flows having $f=1$ (inefficient cooling), here we do not impose any restriction to the flow parameters to start with and let the parameter $f$ to determine self-consistently as the system evolves. Therefore, our model is very general whose special case may be understood as a radiatively inefficient advection dominated flow.  We have explored especially the optically thin flows incorporating bremsstrahlung, synchrotron and inverse Comptonization processes. In Fig. \\ref{opti} we show the variation of the effective optical depth as a function of disc radii for two limiting cases. While flows around rotating black holes appear slightly thinner compared to the corresponding cases of static black holes, in general $\\tau_{\\rm eff}\\lsim 5\\times 10^{-4}$. This verifies our choice of optically thin flows throughout. However, for the present purpose, when the main aim is to understand disk dynamics in the global, viscous, two-temperature regime, we have ignored inverse Comptonization due to bremssstrahlung photons, if any. This may be important in cases of very super-Eddington accretion flows which we plan to explore in future, particularly, in analysing the underlying spectra.  \\begin{figure} \\centering \\includegraphics[width=0.50\\columnwidth,angle=-90]{f22.eps} \\caption{ Variation of the effective optical depth as a function of radial coordinate. Solid ($a=0$) and dotted ($a=0.998$) curves correspond to $M=10$, $\\mdot=100$ and dashed ($a=0$) and dot-dashed ($a=0.998$) curves correspond to $M=10^7$, $\\mdot=0.01$. See Tables 1 and 2 for details. } \\label{opti} \\end{figure}  The temperature of the flow depends on the accretion rate. If the accretion rate is low and thus the flow is radiatively inefficient, then the disc is hot. Such a hot flow is being attempted to model since 1976 (Shapiro, Lightman \\& Eardley 1976) when it was assumed that locally $Q^+\\sim Q^-$ and thus $f\\rightarrow 0$. While the model was successful in explaining observed hard X-rays from Cyg~X-1, it turned out to be thermally unstable. Rees et al. (1982) proposed a hot ion torus model avoiding $f$ to unity. In the similar spirit Narayan \\& Yi (1995) proposed the hot two temperature solution in the assumption of $f\\rightarrow 1$ including the strong advection into the flow. Abramowicz et al. (1995), based on the single temperature model, showed that the optically thin disc flow of accretion rate more than one Eddington does not have an equilibrium solution. However, they did not attempt to solve the complete set of differential equations. Based on some simplistic assumptions they showed the importance of advective cooling. Moreover, a single temperature description does not allow them to include all the underlying physics necessary to describe the cooling processes. In the due course, Mandal \\& Chakrabarti (2005) proposed a two temperature disc solution where the ion temperature could be as high as $\\sim 10^{12}$K. However, they particularly emphasized on how does the shock in the disc flow enable cooling through the synchrotron mechanism, without carrying out a complete analysis of the dynamics. The present paper describes, to our knowledge, the first comprehensive work to model the two temperature accretion flow self-consistently by solving the complete set of underlying equations without any pre-assumptive choice of the flow variables to start with.  The generality lies not only in its construction but also its ability to explain the under-luminous to ultra-luminous sources, stellar mass to supermassive black holes. Table 3 lists the luminosities for a wide range of parameter sets, obtained by our model. It reveals that for a very low mass accretion rate $\\mdot=0.0001$ around a supermassive black hole, the luminosity comes out to be $L\\sim 10^{34}$ erg/sec, which indeed is similar to the observed luminosity from a under-luminous source Sgr~$A^{*}$. In other extreme, for $\\mdot=100$ around a similar black hole, $L\\sim 10^{47}$ erg/sec, similar to what observed from the highly luminous AGNs like PKS~0743-67. On the other hand, when the black hole is considered to be of stellar mass, then at a high $\\mdot=100$, the model reveals $L\\sim 10^{40}$ erg/sec which is similar to the observed luminosity from ULX sources (e.g. SS433).  \\clearpage \\centerline{ Table 3: Luminosity in erg/sec}  \\begin{center} \\begin{tabular}{lllllllllllll} \\hline \\hline $\\mdot$ & $M$ &  $\\gamma$ & $L$ \\\\ \\hline \\hline $0.0001$ & $10^7$       & $1.6$ & $10^{34}$ \\\\ $0.01$ & $10^7$ &  $1.5$ & $10^{38}$  \\\\ $1$  & $10^7$      & $1.35$ & $5\\times 10^{42}$\\\\ $100$  & $10^7$ &  $1.34$ & $10^{47}$\\\\ \\hline \\hline $0.01$  & $10$  & $1.5$ & $10^{33}$\\\\ $1$  & $10$ &  $1.35$ & $7\\times 10^{36}$\\\\ $100$  & $10$   & $1.34$ & $10^{40}$\\\\ \\hline \\hline \\end{tabular} \\end{center}  In general, an increase of accretion rate increases density of the flow which may lead to a high rate of cooling and thus decrease of the cooling factor $f$. Hence, $f$ is higher, close to unity reassembling radiatively inefficient flows, for sub-Eddington accretors, and is lower, sometimes close to zero, for super-Eddington flows. Actual value of $f$ in a flow also depends on the behaviour of hydrodynamic variables which determine the rate of cooling processes. Naturally, as the flow advances from the transition region to the event horizon, $f$ varies between $0$ and $1$. However, if the black hole is considered to be rotating, the flow angular momentum decreases and thus the radial velocity increases. This in turn reduces the residence time of the sub-Keplerian flow hindering cooling processes to complete. This flow is then expected to be hotter and hence $f$ to be higher compared to that around a static black hole. Therefore, the system may tend to be radiatively inefficient, evenif its counter part around a static black hole appears to be an GAAF. However, this also depends on the value of $\\alpha$, as shown in Fig. \\ref{alfacom}. A low value of $\\alpha$ increases the residence time of matter in the disc which helps in cooling processes to complete, rendering a radiatively inefficient flow to switch over to GAAF. This feature may help in understanding the transient X-ray sources.  In all the cases, the ion and electron temperatures merge or tend to merge at around transition radius. This is because, the electrons are in thermal equilibrium with the ions and thus virial around the transition radius, particularly when $\\mdot\\gsim 1$. As the sub-Keplerian flow advances, the ions become hotter and the corresponding temperature increases, rendering the ion-electron Coulomb collisions weaker. The electrons, on the other hand, cool down via processes like bremsstrahlung, synchrotron emissions etc. keeping the electron temperature roughly constant upto very inner disc. This reveals the two temperature flow strictly.  Important point to note is that we have assumed throughout the coupling between the ions and electrons is due to the Coulomb scattering. However, the inclusion of possible nonthermal processes of transferring energy from the ions and electrons (Phinney 1981, Begelman \\& Chiueh 1988) might modify the results. However, as argued by Narayan \\& Yi (1995), the collective mechanism discussed by Begelman \\& Chiueh (1988) may dominate over the Coulomb coupling at either a very low $\\alpha$ or a very low $\\mdot$. Instead, the viscous heating rate of ions is much larger than the collective rate of nonthermal heating of electrons, unless $\\alpha$ is too small what we have not considered in the present cases. Therefore, the assumption to neglect nonthermal heating of electrons is justified.  Now the future jobs should be to understand the radiation emitted by the flows discussed here and to model the corresponding spectra. This will be the ultimate test of the model in order to explain observed data.  \\appendix"
    },
    "0910/0910.1920_arXiv.txt": {
        "abstract": "We explore potential of current and next-generation \\gr\\ telescopes for the detection of weak magnetic fields in the intergalactic medium. We demonstrate that using two complementary techniques, observation of extended emission around point sources and observation of time delays in \\gr\\ flares, one would be able to probe most of the cosmologically and astrophysically interesting part of the  \"magnetic field strength\" vs. \"correlation length\" parameter space. This implies that  \\gr\\ observations with {\\it Fermi} and ground-based Cherenkov telescopes will allow to (a) strongly constrain theories of the origin of magnetic fields in galaxies and galaxy clusters and (b) discover, constrain or rule out the existence of weak primordial magnetic field generated at different stages of evolution of the Early Universe. ",
        "introduction": "Magnetic fields are known to play an important role in the physics of a variety of astrophysical objects, from stars to galaxies and galaxy clusters. The existence of galactic magnetic fields with strengths in the range $1-10\\ \\mu$G is established via observations of Faraday rotation and Zeeman splitting of atomic lines in the radio band and of polarization of starlight in the optical band \\cite{kulsrud,beck08}. Magnetic fields of similar strength are found in the cores of galaxy clusters \\cite{carilli02}. Weaker magnetic fields with strengths in the range $10^{-8}-10^{-7}$~G were recently discovered at the outskirts of galaxy clusters \\cite{xu06,kronberg07}.  Although the strength and spatial structure of magnetic fields in the Milky Way and some other galaxies are reasonably well known today, there is no commonly accepted theory about the origin of these magnetic fields (see \\cite{kronberg94,grasso00,widrow02,kulsrud} for recent reviews of the subject). There is a general agreement that the observed microGauss magnetic fields are the result of amplification of weak \"seed\" fields. The amplification mechanisms under discussion are the so-called \"$\\alpha-\\omega$\" dynamo (in the case of spiral galaxies) and/or compression and turbulent motions of plasma during the galaxy/cluster formation processes.  The nature of the initial weak seed fields for the dynamo or turbulent amplification is largely unknown \\cite{kronberg94,grasso00,widrow02,kulsrud}. It might be that the seed fields are produced during the epoch of galaxy formation by electrical currents generated by the plasma experiencing gravitational collapse within a proto-galaxy \\cite{pudritz89,gnedin00}, or ejected by the first supernovae \\cite{rees87} or active galactic nuclei \\cite{donnert09}. Otherwise, the seed fields might originate from still earlier epochs of the Universe expansion, down to the cosmological phase transitions or inflation times \\cite{grasso00}.  Wide uncertainties in both the mechanism of amplification of the seed fields and in the nature of the seed fields themselves have led, over the last half-a-century, to the appearance of a long-standing problem of the \"origin of cosmic magnetic fields\" (in galaxies and galaxy clusters).  It is clear that the clue for the solution of this problem might be given by the measurements of the initial seed fields. However, up to recently there was little hope that the extremely weak fields  outside galaxies and galaxy clusters would ever be detected.  In what follows we show that direct measurements of the seed fields and derivation of constraints on their nature become possible with the newly available observations in the very-high-energy \\gr\\ band with space and ground-based \\gr\\ telescopes such as {\\it Fermi}, HESS, MAGIC, VERITAS and, in the near future, CTA, AGIS and HAWC. The method of measurement of \"ExtraGalactic\" Magnetic Fields (EGMF) with \\gr\\ telescopes is based on the possibility of detection of emission from electromagnetic cascade initiated by the primary \\gr s emitted by an extragalactic source and developing throughout the InterGalactic Medium (IGM) along the line of sight toward the source \\cite{plaga,neronov07,japanese,elyiv09,kachelriess09}. Based on the knowledge of sensitivity of existing and future \\gr\\ telescopes, we find the range of EGMF parameters, such as the field strength $B$ and the correlation length $\\lambda_B$, in which the EGMF is accessible for the measurements with one of the two available measurement techniques (imaging \\cite{neronov07,elyiv09,kachelriess09} or timing \\cite{plaga,japanese} of the cascade signal). We demonstrate that most of the astrophysically and cosmologically interesting range of EGMF parameters could be probed with \\gr\\ observations.  The plan of the paper is as follows. In Section \\ref{sec:experiment} we summarize the existing bounds on the strength and correlation length of EGMF which come mostly from radio observations. In Section \\ref{sec:cosmology} we discuss limits on the cosmological magnetic fields from cosmology. Then, in Section \\ref{sec:theory} we compare the existing bounds to the theoretical predictions of two classes of models (\"astrophysical\" vs. \"cosmological\" models)  of the \"seed\" fields and show that model predictions normally fall largely below the existing bounds. In Sections \\ref{sec:gammaray} -- \\ref{sec:strongB} we first summarize the methods of measurement of EGMF with \\gr\\ telescopes and then estimate the ranges of EGMF parameters which can be probed with different observational techniques and different telescopes. Finally, in Section \\ref{sec:conclusions} we draw conclusions from our study. ",
        "conclusions": "\\label{sec:conclusions}  We have discussed the prospects of detection of weak magnetic fields in the intergalactic medium with the novel techniques of timing and imaging observations with ground and space-based \\gr\\ telescopes. These techniques enable measurements of extremely weak magnetic fields with strengths much lower than the ones accessible for the measurements with radio telescopes (via Zeeman splitting and/or Faraday rotation techniques).  We have demonstrated that using \\gr\\ observations one can detect, or rule out the possibility of existence of cosmologically or astrophysically produced \"seed\" magnetic fields in the voids of the large scale structure, which are conjectured to exist in a range of theories of the origin of magnetic fields in galaxies and galaxy clusters (see Fig. \\ref{fig:exclusion_theory}).  To summarize, we find that discovery or non-detection of weak magnetic fields  in the voids of the large scale structure should provide, in the nearest future, a decisive test of the theories of the origin of cosmic magnetic fields."
    },
    "0910/0910.1934_arXiv.txt": {
        "abstract": "The dynamical properties of a model of dark energy in which two scalar fields are coupled by a noncanonical kinetic term are studied. We show that overall the addition of the coupling has only minor effects on the dynamics of the two-field system for both potentials studied, even preserving many of the features of the assisted quintessence scenario. The coupling of the kinetic terms enlarges the regions of stability of the critical points. When the potential is of an additive form, we find the kinetic coupling has an interesting effect on the dynamics of the fields as they approach the inflationary attractor, with the result that the combined equation of state of the scalar fields can approach $-1$ during the transition from a matter dominated universe to the recent period of acceleration. ",
        "introduction": "\\label{sec:Introduction} One goal of cosmology is to understand the origin of the observed accelerated expansion of the universe (see \\cite{Copeland:2006wr} and references therein). To date, there are several suggestions, including the cosmological constant, slowly evolving scalar fields and modifications to Einstein's theory of general relativity. Since scalar fields are predicted by many particle physics theories, scalar field models of dark energy, such as quintessence \\cite{Wetterich:1987fm, Ratra:1987rm, Martin:2008qp} or k-essence \\cite{Chiba:1999ka, ArmendarizPicon:2000dh, ArmendarizPicon:2000ah} have been studied in considerable depth in the past. One requirement a satisfactory model of dark energy must fulfill is that it leads to an equation of state (EOS) close to $w=-1$, in order to agree with current observational data. For single scalar fields, exponential potentials with slope $\\lambda$ lead to a scalar field dominated universe with late-time accelerated expansion if $\\lambda<\\sqrt{2}$. In this case, the duration of the matter dominated epoch depends on the initial conditions for the scalar field. Thus, the situation is not better than that with a cosmological constant. If, on the other hand, $\\lambda>\\sqrt{3(1+w)}$, the scalar field scales with the dominant fluid (with EOS $w$ \\cite{Ferreira:1997hj, Copeland:1997et}). This would help the initial condition problem, but unfortunately, the scaling solution and the accelerating solution are mutually exclusive. The situation with inverse power-law potentials is better in the sense that a wide range of initial conditions end up with the same cosmology at late times, but in order for the theory to be consistent with observational data, the exponent has to be small.  On the other hand, if quintessence is indeed the scenario realized in nature, the quintessence sector might turn out to be rather nontrivial. There might, for example, be several scalar fields interacting \\cite{Barreiro:1999zs,Coley:1999mj,Blais:2004vt,Kim:2005ne} and/or the scalar fields might interact with matter in the universe \\cite{Amendola:1999er,TocchiniValentini:2001ty, Brookfield:2007au, Bean:2008ac}. In this paper we will study interacting scalar fields as a model for dark energy and study whether this can shed some light on the issues in quintessence model building. Models with multiple scalar fields have been considered in the past: partly to study isocurvature (entropy) perturbations (e.g. \\cite{Langlois:1999dw,Gordon:2000hv,Byrnes:2006fr,Langlois:2008vk}) and/or non-Gaussianity generated during inflation (see e.g. \\cite{Alishahiha:2004eh,Rigopoulos:2005ae,Hattori:2005ac,Cai:2009hw} and references therein). It was also found that the cumulative effect of many scalar fields could relax the constraints on the inflationary potential \\cite{Liddle:1998jc,Malik:1998gy,Calcagni:2007sb}. These ideas have also been used in models of dark energy (assisted quintessence  \\cite{Kim:2005ne,Tsujikawa:2006mw,Ohashi:2009xw}). Like the case with inflation in the very early universe, there is no reason why dark energy is not driven by several interacting scalar fields. In a scenario with several scalar fields one might hope to find an explanation for the coincidence problem, i.e. why dark energy dominates today and not earlier during the cosmic history. In single scalar field models this is rather difficult to achieve (see e.g. \\cite{Copeland:2006wr}).  We consider a simple extension of the standard case with canonical normalized fields by allowing for a cross-term in the kinetic energy of the scalar field, which, in the case of a homogeneous and isotropic universe, is proportional to $\\dot\\phi\\dot\\chi$. This is the simplest extension possible in the two-field case without changing the properties of the potential energy. Kinetic interactions between multiple scalar fields have been previously studied in the context of models in which the dark energy equation of state can become less than $-1$ \\cite{Sur:2009jg,Chimento:2008ws}.  The paper is organized as follows: in the next section we present the model and discuss some aspects of the dynamics of the system, such as effective exponents. In Secs. \\ref{sec:assist} and \\ref{sec:soft} we perform a critical point analysis for two types of exponential potentials and  in Sec. \\ref{sec:assist_num} we present the results of a numerical analysis. In Sec. \\ref{sec:var_a} we consider the case of a varying coupling function. Our conclusions can be found in Sec. \\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion} Although the use of scalar fields to model the acceleration of the universe could well be described as standard practice in cosmology, the fundamental physics that underpins these phenomenological models is by no means clear. It is wise, therefore, to consider the possibility that the observed acceleration of the universe may have more exotic origins. The recent focus on scalar fields with noncanonical kinetic terms derived from string-inspired models has shown that models that deviate from the standard lore are capable of producing fruitful results that display qualitatively different behaviour. In this work we have considered a model of dark energy in which the kinetic terms of two scalar fields are coupled by a term in the Lagrangian of the form $a(\\p_\\alpha\\phi\\p^\\alpha\\chi)$. This allowed us to study the simplest extension of the standard case without modifying the potential energy. After analysing the phase plane structure of the model for two different potentials, we found that the basic features of the uncoupled case were preserved and the most important effect of the extra degree of freedom was to change the range of parameters that lead to the different types of solutions. In particular, there exist stable points corresponding to scalar field-dominant or scaling solutions for every value of $\\lambda$, $\\mu$ and $a$ considered.  In the case of the assisted potential we found that for negative $a$, critical points C1, D1, C2 and D2 that existed but were unstable in the case without the kinetic coupling became stable. This means that when $a$ goes from positive to negative values, there are three different regions in the $\\lambda,\\mu$-parameter space that give rise to accelerated solutions, each with a different value of the EOS, $w_{\\rm fields}$. This is less interesting from a phenomenological point of view, however, as except when the values of the potential exponents are very small, negative $a$ values have $w_{\\rm fields}$ larger than the observed dark energy EOS. In the case of positive $a$, one finds the EOS to be closer to $-1$ than that in the model with $a=0$, though in the region of interest ($w_{\\rm fields}\\approx-1$) the effect of the kinetic coupling is extremely slight. The soft potential evinced a similar dependence on $a$, though it is of less interest as a cosmological model  as accelerated behaviour only occurs with small values of the exponents in all cases.  Although models with exponential potentials exhibit both scaling and scalar field-dominant solutions, it is difficult to solve the coincidence problem without fine-tuning the initial conditions or allowing the parameters to vary in time. In Sec. \\ref{sec:assist_num} we discuss how the kinetic coupling affects the dynamics of the fields as they evolve towards the critical points. When $a$ is close to $2$ one finds that the value of $w_{\\rm fields}$ becomes very close to $-1$ for a brief period before the scalar fields dominate. This is a partial resolution of the perennial problem of the EOS being too large.  This analysis is revealing as the addition of a cross-kinetic term has only a relatively minor effect on the dynamics of the two-field system for both cases studied, even preserving many of the features of the assisted quintessence scenario. In the case of the assisted potential, the kinetic coupling has an interesting effect on the dynamics of the fields as they approach the stable solution, with the result that the EOS of the scalar fields can approach $-1$ during the transition from a matter dominated universe to the recent period of acceleration, something that does not occur in assisted quintessence or the single field case. A natural extension to this work is to consider the case when the scalar fields are coupled directly to the matter fluid by a similar mechanism to that described in \\cite{TocchiniValentini:2001ty}. In this case, one would expect the value and stability ranges of the critical points to change dramatically due to the extra degree(s) of freedom, but it would be interesting to see whether the additive potential still allows large values of the potential exponents. The work could also be extended by treating $\\lambda$ and $\\mu$ as dynamical variables in a manner similar to \\cite{Ng:2001hs}, in which case it may be possible to obtain solutions with a large range of initial conditions that track the background solution and later dominate.  As we have discussed, our results are also applicable in the case of a slowly varying coupling function $a(\\phi,\\chi)$  as well as the case in which the coupling undergoes a rapid change but is constant (or evolving very slowly) after the transition."
    },
    "0910/0910.3877_arXiv.txt": {
        "abstract": "We present the near-infrared luminosity and stellar mass functions of galaxies in the core of the Shapley supercluster at $z$=0.048, based on new $K$-band observations carried out at the United Kingdom Infra-Red Telescope with the Wide Field Infrared Camera in conjunction with $B$- and $R$-band photometry from the Shapley Optical Survey, and including a subsample ($\\sim$650 galaxies) of spectroscopically confirmed supercluster members. These data sets allow us to investigate the supercluster galaxy population down to M$_K^\\star$+6 and $\\mathcal{M}$=10$^{8.75}$M$_\\odot$. For the overall 3\\,deg$^2$ field the $K$-band luminosity function (LF) is described by a Schechter function with M$_K^\\star$=--24.96$\\pm$0.10 and $\\alpha$=--1.42$\\pm$0.03, a significantly steeper faint-end slope than that observed in field regions. We investigate the effect of environment by deriving the LF in three regions selected according to the local galaxy density, and observe a significant ($2\\sigma$) increase in the faint-end slope going from the high- ($\\alpha$=--1.33) to the low-density ($\\alpha$=--1.49) environments, while a faint-end upturn at M$_{K}>$--21 becomes increasingly apparent in the lower density regions.  The galaxy stellar mass function (SMF) is fitted well by a Schechter function with log$_{10}$($\\mathcal{M}^\\star$)=11.16$\\pm$0.04 and $\\alpha$=--1.20$\\pm0.02$. The SMF of supercluster galaxies is also characterised by an excess of massive galaxies that are associated to the brightest cluster galaxies. While the value of $\\mathcal{M}^*$ depends on environment increasing by 0.2 dex from low- to high-density regions, the slope of the galaxy SMF does not vary with the environment. By comparing our findings with cosmological simulations, we conclude that the environmental dependences of the LF are not primary due to variations in the merging histories, but to processes which are not treated in the semi-analytical models, such as tidal stripping or harassment. In field regions the SMF shows a sharp upturn below $\\mathcal{M}$=10$^{9}$M$_\\odot$, close to our mass limit, suggesting that the upturns seen in our $K$-band LFs, but not in the SMF, are due to this dwarf population. The environmental variations seen in the faint-end of the $K$-band LF suggests that these dwarf galaxies, which are easier to strip than their more massive counterparts, are affected by tidal/gas stripping upon entering the supercluster environment. ",
        "introduction": "\\label{intro} The properties and evolution of galaxies are strongly related to their environment (e.g. Blanton et al. \\citeyear{bla05a}; Rines et al. \\citeyear{rin05}; Baldry et al. \\citeyear{BBB06}), through the mass and merging histories of their host dark matter halos, and the impact of different physical mechanisms (e.g. Treu et al. \\citeyear{tre03}) that are linked in various ways to the local galaxy density and the properties of the intergalactic medium.  In the local Universe this environmental dependence has been investigated and observed in the distribution of galaxy luminosities and stellar masses, providing constraints on the assembly of galaxies over cosmic time (see below).  Since the NIR light is dominated by established old stellar populations rather than by recent star-formation activity, the NIR LF can be considered as reliable estimator of the stellar mass function (SMF, Gavazzi et al. \\citeyear{GPB96}; Bell \\& de Jong \\citeyear{BdJ01}) and the shape of the NIR LF constrains the scenarios of galaxy formation (e.g. Benson et al. \\citeyear{BBF03}), with different energetic feedback processes from supernovae and AGN required to simultaneously fit the LFs at the faint and bright ends respectively \\citep{BBM06}. Early versions of the hierarchical galaxy formation model predicted a decrease of the abundance of massive galaxies with redshift (Kauffmann \\& Charlot \\citeyear{KC98}) which have continued forming until recent times through processes of merging and accretion, and a steep mass function due to the presence of a large number of faint dwarf galaxies witnessing the small dark halo formation in the early universe (e.g. Kauffmann et al. \\citeyear{KWG93}).  The NIR LFs observed at different redshifts turn out to be well described by Schechter functions, although there is some evidence for an excess of bright galaxies with respect to the best fitting Schechter function (e.g. Jones et al. \\citeyear{JPC06}), while a faint-end upturn has also been observed in some cases (e.g. De Propris \\& Pritchet \\citeyear{DPP98}; Balogh et al. \\citeyear{BCZ01}; Jenkins et al. \\citeyear{JHM07}).  The results supporting the hierarchical scenario (e.g. Kauffmann \\& Charlot \\citeyear{KC98}) are contrasted by other works that bring it into question (Kodama \\& Bower \\citeyear{KB03}; De Propris et al. \\citeyear{DSE99}; \\citeyear{DSE07}) or cannot interpret univocally their findings (e.g. Drory et al. \\citeyear{DBF03}). For instance, there is a consensus that the evolution of the characteristic absolute luminosity M$^\\star$ for both field and cluster galaxies can be described by a passively evolving population formed in a single burst at redshift $z$=1.5-2 (e.g. Lin et al. \\citeyear{LMG06}; Wake et al. \\citeyear{WNE06}; De Propris et al. \\citeyear{DSE99}; \\citeyear{DSE07}). On the other hand, Drory et al. (\\citeyear{DBF03}) attribute the observed evolution in the $K$-band LF of cluster galaxies, either to a change in the mass-to-light ratio alone (i.e. passive evolution), or to a combination of changes in M/L and stellar mass which could be due to star formation and/or to merging or accretion. Recent semi-analytic models incorporating AGN feedback have been able to reproduce well the evolution of the SMF over $0{<}z{<}5$ (Bower et al. \\citeyear{BBM06}), predicting the observed population of massive galaxies at $z{>}2$. However, in order to properly use the NIR, LF to disentangle between possible scenarios of galaxy evolution, one has to consider the dependence of the LF on environment, galaxy colour, spectral type and morphology.  The LF of field galaxies at NIR wavebands has been primarily investigated through the Two Micron All Sky Survey\\footnote{http://www.ipac.caltech.edu/2mass/} (2MASS) often complemented by either optical photometry from the Sloan Digital Sky Survey\\footnote{http://www.sdss.org/} (SDSS) or wide-field spectroscopic surveys such as the 2dF Galaxy Redshift Survey\\footnote{http://www.mso.anu.edu.au/2dFGRS/} (2dFGRS). These works (e.g. Kochanek et al. \\citeyear{KPF01}; Cole et al. \\citeyear{CNB01}) agree in describing the NIR LFs with Schechter functions characterised by a rather flat faint-end slope $\\alpha$ from --0.77 to --0.96 and which are independent of the morphological type (Kochanek et al. \\citeyear{KPF01}), in contrast with that found in optical surveys where the faint-end slope is steeper for late-type galaxies. On the contrary, Bell et al. (\\citeyear{BMK03}) found that NIR LF has a brighter characteristic luminosity and shallower slope for early-type galaxies with respect to the later types. More recently, Jones et al. (\\citeyear{JPC06}), using the 6dF Galaxy Redshift Survey\\footnote{http://www.aao.gov.au/local/www/6df/} (6dFGRS, Jones et al. \\citeyear{JSC04}) derived NIR LFs for field galaxies 1-2\\,mag deeper in absolute magnitude with respect to the previous works. They found that the Schechter function is not ideal to reproduce the data since it cannot match the bright- and faint-end simultaneously, due to an excess of galaxies at magnitudes brighter than M$^\\star$.  \\begin{figure*} \\centerline{\\resizebox{17.5cm}{!}{\\includegraphics{fig1.eps}}} \\caption{The Shapley NIR survey (thick red) is shown superimposed to the surface density of $\\mathrm{R}<21.0$ galaxies as derived from the SOS (see text). Isodensity contours are shown at intervals of 0.25 galaxies arcmin$^{-2}$, with the thick contours corresponding to 0.5, 1.0 and 1.5 galaxies arcmin$^{-2}$, the densities used to separate the three cluster environments. The multi-wavelength photometric coverage of ACCESS on the SSC are also shown. Magenta: 24\\,$\\mu$m, green: 70\\,$\\mu$m, blue dashed: 150\\,nm and 250\\,nm.} \\label{fig1} \\end{figure*}  Balogh et al. (\\citeyear{BCZ01}) investigated the dependence of the infrared galaxy luminosity function and the associated galaxy SMF on environment and spectral type by means of 2MASS and Las Campanas Redshift Survey (LCRS, Shectman et al. \\citeyear{SLO96}) for galaxies brighter that M$_J$=--19\\,mag. In field environments the LF of galaxies with emission lines turns out to have a much steeper faint-end slope ($\\alpha$=--1.39) compared to that of galaxies without emission lines ($\\alpha$=--0.59). On the other hand, in the cluster environment, even the non-emission line galaxies have a steep faint-end LF ($\\alpha$=--1.22). This difference is almost entirely due to the non-emission line galaxies which dominate the cluster population, and present a slope close to that of the overall field. Thus, they suggested that the cluster population is built up by accreting field galaxies with little effect other than the cessation of star formation. Differences in the shape of the LF for late- and early-type cluster galaxies has been found by Huang et al. (\\citeyear{HGC03}): the late-type galaxies having a systematically fainter M* and steeper faint-end slope. A possible faint-end upturn in the $H$-band LF was already suggested by De Propris et al. (\\citeyear{DES98}, see also Andreon \\& Pell\\'o \\citeyear{AP00}) for the Coma cluster outlining the steep trend of dwarf galaxies down to M*+5, even if they did not provide a precise estimate of the faint-end slope because of possible field contaminations. An increase of the faint-end slope of Coma was recently observed at 3.6\\,$\\mu$m by Jenkins et al. (\\citeyear{JHM07}) indicating a large number of faint red galaxies. However, Rines \\& Geller (\\citeyear{rin08}) found no such upturn for Virgo, based on a fully spectroscopically confirmed sample, and suggested that many of the photometrically-selected red sequence galaxies which contribute to the upturns seen in other clusters are background galaxies.  Finally, the tight correlation between the total galaxy NIR luminosity and the cluster binding mass (Lin et al. \\citeyear{LMS03}, \\citeyear{LMS04}; Ramella et al. \\citeyear{RBG04}) allows to probe that the global cluster $K$-band mass-to-light ratio decreases with cluster radius (Rines et al. \\citeyear{RGD04}) showing that the environment affects the shape of the LF also within the clusters.  We note that most of the previous works are based on the 2MASS data which has a detection sensitivity (10$\\sigma$) of $K$=13.1\\,mag for extended sources (Cole et al. \\citeyear{CNB01}, $K$=13.57\\,mag according to Bell et al. \\citeyear{BMK03}), limiting studies of the environmental impact on the NIR LF to only a sample of local clusters (e.g. Virgo), or limiting to magnitudes $<$M$^\\star$+2 (e.g. Rines et al. \\citeyear{RGD04}). However, the dominant processes that quench star-formation, and therefore transform galaxies, depend crucially on the galaxy mass (e.g. Haines et al.  \\citeyear{HLM06}; \\citeyear{HGL07}), and the strong bimodality in the properties of galaxies about a characteristic stellar mass of $\\sim 3\\times 10^{10} M_{\\odot}$ ($\\sim$ M*+1, Kauffmann et al.  \\citeyear{KHW03}) implies fundamental differences in the formation and evolution of giant and dwarf galaxies (e.g. Dekel \\& Birnboim \\citeyear{DB06}; Kere\\^s et al. \\citeyear{KKW05}). This issue needs data-sets reaching much fainter luminosities than those of 2MASS to be investigated in order to obtain, in general, an overall picture of galaxy evolution and, in particular, to establish the contribution of the dwarf galaxy population to the total stellar mass in the local universe and the physical origin of the claimed faint-end upturn. The recent development of wide-field NIR imagers on 4-m class telescopes such as UKIRT/WFCAM and KPNO/NEWFIRM has opened the possibility of NIR surveys to be $>100\\times$ more sensitive covering many square degrees (e.g. UKIDSS), allowing the $K$-band LF of nearby clusters to be obtained covering not only the cluster cores, but the entire virialized regions.  In this context we study the $K$-band LF of the Shapley supercluster core (SSC) down to the dwarf regime (reaching $\\sim$M$^\\star$+6) with the aim of i) quantifying the environmental impact on the shape of the NIR LF; ii) deriving the stellar masses of the supercluster galaxies; iii) investigating the mechanisms driving galaxy evolution as function of galaxy mass.  This work is carried out in the framework of the joint research programme ACCESS aimed at determining the importance of cluster assembly processes in driving the evolution of galaxies as a function of galaxy mass and environment within the Shapley supercluster (see Sect.~\\ref{ACCESS}). In Sect.~\\ref{sec:3}, we describe the data-sets.  In Sect.~\\ref{sec:4}, we derive the NIR galaxy luminosity functions obtained through background subtraction in the whole observed field and we study the ongoing effects of environment by comparing the LFs of galaxies in three different regions of the supercluster, characterized by high-, intermediate- and low-density, we also compare NIR and optical LFs. The galaxy stellar mass function is presented in Sect.~\\ref{sec:5}.  The results are discussed in Sect.~\\ref{sec:6} and the summary and conclusions of this work is given in Sect.~\\ref{sec:7}.  Throughout the paper we adopt a cosmology with $\\Omega_M$=0.3, $\\Omega_\\Lambda$= 0.7, and H$_0$=70 km s$^{-1}$Mpc$^{-1}$. According to this cosmology 1 arcmin corresponds to 60 kpc at $z$=0.048 and the distance modulus is 36.66. ",
        "conclusions": "\\label{sec:7}  It is well known that the NIR luminosities provide a more reliable estimate of the stellar masses compared to the optical ones due to the fact that the mass-to-light ratios in the NIR vary by at most a factor 2 across a wide range of star formation histories (Bell \\& de Jong \\citeyear{BdJ01}) in comparison to the much larger variations of the M/L ratios (up to a factor 10) observed at optical wavelengths. In addition the effects of the extinction are much weaker at infrared wavelengths that in the optical ones, and k-corrections for infrared colours are only weakly dependent on the Hubble type and vary slowly with redshift (e.g. Poggianti \\citeyear{pog97}).  In our study we exploit new deep ($K$=18\\,mag) $K$-band imaging of the SSC complemented by the deep optical imaging down to $B$=22.5\\,mag and $R$=22\\,mag from the SOS, and a spectroscopically confirmed supercluster sample of $\\sim$650 galaxies across the same field which is $\\sim$90\\% complete at $R<$16\\,mag. We present an analysis of the $K$-band LF of galaxies as a function environment and we derive the galaxy SMF in order to constrain the physical mechanisms that transform galaxies in different environments as function of galaxy mass.  Our results are summarized as follows.  \\begin{description}  \\item[-] The $K$-band LF can be fitted by a single Schechter function, with $M_{K}^{*}=-24.96\\pm0.10$ and $\\alpha=-1.42\\pm0.03$ in agreement with previous works of comparable depth.  \\item[-] The $K$-band LF faint-end slope becomes steeper from high- to low-density environments varying from --1.33 to --1.49, being inconsistent at the 2$\\sigma$ c.l. (see Tab~\\ref{fitsLF}), indicating that the faint galaxy population increases in low-density environments. Such an environmental dependence confirms our finding for the optical LFs derived in the same supercluster regions although the changes in slope are less dramatic at NIR wavebands.  \\item[-] The observed trend of the galaxy SMF presents a slope in agreement with previous works concerning the SMF of field galaxies (Cole et al. \\citeyear{CNB01}; Panter et al. \\citeyear{pan04}).  The value of $\\mathcal{M}^\\star$ is higher with respect to that observed in the field (e.g. Bell et al. \\citeyear{BMK03}), as expected for cluster environment. The SMF of supercluster galaxies is characterised by an excess of massive galaxies that is associated to the cluster BCGs.  Discrepancies with previous work that observed a strong faint-end upturn (Baldry et al. \\citeyear{BGD08} and Jenkins et al. \\citeyear{JHM07}) can be related to the different mass ranges investigated and/or environmental differences in the analysed structures.  \\item[-] Differently from the LF no environment effect is found in the slope of the SMFs. On the other hand, the $\\mathcal{M}^\\star$ increase from low- to high-density regions and the excess of galaxies at the bright-end is also dependent on the environment. This trend is in general agreement with the results of Baldry et al. (\\citeyear{BBB06}).  \\end{description}  In order to interpret our findings, we use the Millennium simulation which produce DM halo and galaxy catalogues based on SAMs. The cluster NIR LFs obtained using the simulated catalogues do not show any significant variation with cluster-centric radius, thus suggesting that the variations observed in the LFs of the Shapley supercluster are not driven by variations in the merging histories of the galaxies, but are likely related to processes such as tidal stripping or harassment of infalling galaxies.  By comparing the effect of environment at optical and NIR wavebands in shaping the LFs and taking into account that the slope of the galaxy SMF is invariant with respect to the environment, it seems that the physical mechanism responsible for the transformation of galaxies properties in different environment are related to the quenching of the star formation rather than mass-loss.  This suggests that the mechanism responsible for shaping the LF and SMF is partially related to the mass loss due to tidal stripping or galaxy harassment, but gas stripping by the ICM can also affect the galaxy population removing the cold gas supply and thus rapidly terminating ongoing star-formation. These mechanisms all require a dense ICM and so their evolutionary effects on massive galaxies are limited to the cores of clusters, but can extend to poorer environments for dwarf galaxies which are easier to strip. The infalling galaxies are probably late-type (see MMH06) that are affected by gas stripping entering in the supercluster.  On the other hand, the depth of the NIR survey could affect the present results by not allowing us to detect an upturn of the SMF which can be detected at a mass range lower than reached here (see Jenkins et al. \\citeyear{JHM07}). In order to identify the mechanisms which drive galaxy evolution in the supercluster environment we are undertaking a survey with the integral field spectrograph WiFeS which will provide a unique data set to investigate in details the stellar populations and kinematics for a subsample of the Shapley galaxies."
    },
    "0910/0910.3793_arXiv.txt": {
        "abstract": "We examine deep \\emph{XMM-Newton} Reflection Grating Spectrometer (RGS) spectra from the cores of three X-ray bright cool core galaxy clusters, Abell~262, Abell~3581 and HCG~62. Each of the RGS spectra show Fe~\\textsc{xvii} emission lines indicating the presence of gas around 0.5~keV. There is no evidence for O~\\textsc{vii} emission which would imply gas at still cooler temperatures. The range in detected gas temperature in these objects is a factor of 3.7, 5.6 and 2 for Abell~262, Abell~3581 and HCG~62, respectively. The coolest detected gas only has a volume filling fraction of 6 and 3 per~cent for Abell 262 and Abell 3581, but is likely to be volume filling in HCG~62. \\emph{Chandra} spatially resolved spectroscopy confirms the low volume filling fractions of the cool gas in Abell~262 and Abell~3581, indicating this cool gas exists as cold blobs. Any volume heating mechanism aiming to prevent cooling would overheat the surroundings of the cool gas by a factor of 4. If the gas is radiatively cooling below 0.5~keV, it is cooling at a rate at least an order of magnitude below that at higher temperatures in Abell~262 and Abell~3581 and two-orders of magnitude lower in HCG~62. The gas may be cooling non-radiatively through mixing in these cool blobs, where the energy released by cooling is emitted in the infrared. We find very good agreement between smooth particle inference modelling of the cluster and conventional spectral fitting. Comparing the temperature distribution from this analysis with that expected in a cooling flow, there appears to be a even larger break below 0.5 keV as compared with previous empirical descriptions of the deviations of cooling flow models. ",
        "introduction": "The hot (few $10^7$K) intracluster medium (ICM) in galaxy clusters is primarily seen in emission in the X-ray waveband. The ICM contains most of the baryonic cluster mass. In a large fraction of clusters of galaxies, observations show that the X-ray surface brightness steeply rises towards their centres (e.g. \\citealt{Stewart84}). In addition, spectrally-derived temperature profiles of clusters typically show drops in temperatures by a factor of 2 or 3 from the outskirts (e.g. \\citealt{AllenSchmidtFabian01}). As the X-ray surface brightness is proportional to the integral of the density squared along a line of sight, peaked surface brightness profiles imply that the gas in the cores is cooling much more rapidly than in the outskirts.  In the centre, the mean radiative cooling time, often calculated as the ratio of the luminosity of a region to its enthalpy, often drops below 1~Gyr.  Assuming steady-state and the absence of heating, there will be material with short cooling times in the cluster centre cooling out of the X-ray emitting band. As the volume of this material becomes much smaller, there must be a flow of material in to replace it to preserve the steady state. The steep surface brightness profiles imply rates of 10s to 1000s of solar masses per year cooling out of the X-ray band (see \\citealt{Fabian94} for a review). This cooling material would be expected to eventually give rise to star formation.  This picture changed when high spectral resolution X-ray studies of nearby clusters of galaxies using the Reflection Grating Spectrometer (RGS) instruments on \\emph{XMM-Newton} revealed a lack of cool X-ray emitting gas in these objects \\citep{Tamura01b, Tamura01a, Peterson01, Kaastra01, Sakelliou02, Peterson03}. The main spectral indicator missing in these spectra are the emission lines of Fe~\\textsc{xvii}, which are strong from material between 0.15 and 0.8~keV (1.7 to 9.4~MK).  The general accepted picture is that there is a lack of material below a factor 2 or 3 of the outer temperature, which matches the results from \\emph{Chandra} spatially-resolved temperature profiles. This is not completely correct, however. Deep observations of nearby clusters show a much larger range of temperature. Fe~\\textsc{xvii} emission lines have been seen in Centaurus \\citep{SandersRGS08}, showing a temperature range of more than 10, Abell 2204 \\citep{SandersA220409}, showing a range of 15, 2A~0335+096 \\citep{Sanders2A033509}, showing a range of 8. Fe~\\textsc{xvii} emission was also detected in Abell~262, M87 and NGC~533 in a large study by \\cite{Peterson03}. There is much less material than would be expected to be seen in the case of steady-state radiative cooling without any heating, however.  There is evidence that the central active galactic nuclei (AGN) in these objects can energetically prevent much of the cooling (see reviews by \\citealt{PetersonFabian06} and \\citealt{McNamaraNulsen07}). This heating may be by the inflation of cavities (radio lobes) by the AGN or the subsequent dissipation of sound waves generated by the inflation. Such cavities are found in almost all nearby clusters which require heating \\citep{DunnFabian06}.  If AGNs are responsible for preventing cooling in cluster cores, there are still remaining issues. In Centaurus the cooling time of the lowest detected component is only $10^7$~yr \\citep{SandersRGS08}. If no cooling is taking place, then feedback must be able to operate on these timescales. Feedback must also be able to operate over the much longer cluster lifetime. Star formation in Centaurus must have occured slowly over the last 8~Gyr, or all must have been at earlier times \\citep{SandersEnrich06}.  We assume a Hubble constant of $70 \\kmpspMpc$ and use the relative Solar metallicities of \\cite{AndersGrevesse89}.  \\begin{table*} \\caption{Details of the targets and individual \\emph{XMM-Newton} observations. The exposure times given are for the RGS1 instrument after cleaning. The absorption quoted is the Galactic absorption from \\protect\\cite{Kalberla05}. References in superscript for redshifts are (1) \\protect\\cite{StrubleRood99}, (2) \\protect\\cite{Johnstone98} and (3) \\protect\\cite{ZabludoffMulchaey00}.} \\begin{tabular}{llllll} \\hline Cluster      & Redshift & Absorption  & Observation & Date       & Exposure \\\\ &    & ($10^{20}\\psqcm$) &             &            & (ks)\\\\ \\hline Abell 262    &$0.0163^1$& 5.67 & 0109980101  & 2001-01-16 & 26.2  \\\\ &          &      & 0504780101  & 2007-07-12 & 121.2 \\\\ &          &      & 0504780201  & 2007-07-18 & 41.4  \\\\ &          &      & total       &            & 188.8 \\\\ Abell 3581   &$0.0218^2$& 4.36 & 0205990101  & 2004-01-29 & 43.5  \\\\ &          &      & 0504780301  & 2007-08-01 & 116.3 \\\\ &          &      & 0504780401  & 2007-08-03 & 27.7  \\\\ &          &      & total       &            & 187.5 \\\\ HCG 62       &$0.01453^3$&3.31 & 0112270701  & 2003-01-15 & 12.4  \\\\ &          &      & 0504780501  & 2007-06-26 & 110.1 \\\\ &          &      & 0504780601  & 2007-06-29 & 33.6  \\\\ &          &      & total       &            & 156.6 \\\\ \\hline \\end{tabular} \\label{tab:sample} \\end{table*}  \\begin{figure*} \\includegraphics[width=0.8\\textwidth]{fig01.pdf} \\caption{\\emph{Chandra} images of the cores of the three objects examined in this paper. Images have been smoothed with a 1 arcsec Gaussian. These images show the \\emph{Chandra} observations in Table~\\ref{tab:chandraobs}.} \\label{fig:chandraimages} \\end{figure*}  \\begin{figure*} \\centering \\includegraphics[width=0.7\\textwidth]{fig02.pdf} \\caption{(Top left panel) Deprojected densities measured from \\emph{Chandra} spectral fitting (shown as points) and from surface brightness deprojection (shown as solid lines). (Top right panel) Deprojected temperature measured from \\emph{Chandra} spectral fitting (shown as points) and surface brightness deprojection (shown as lines). (Bottom left panel) Cooling time profiles measured from surface brightness deprojection (shown as small points) and RGS (shown as large points). (Bottom right panel) Cumulative mass deposition rates derived from \\emph{Chandra} surface brightness deprojection.} \\label{fig:chandraprofiles} \\end{figure*}  \\begin{table} \\caption{\\emph{Chandra} datasets analysed.} \\begin{tabular}{llll} \\hline Target       & Obs-ID & Date       & Clean exposure (ks)\\\\ \\hline Abell 262    & 7921   & 2006-11-20 & 110.7 \\\\ Abell 3581   & 1650   & 2001-06-07 & 7.2 \\\\ HCG 62       & 921    & 2000-01-25 & 48.4 \\\\ \\hline \\end{tabular} \\label{tab:chandraobs} \\end{table} ",
        "conclusions": ""
    },
    "0910/0910.5726_arXiv.txt": {
        "abstract": "We analyze the $>$100 MeV data for 3 GRBs detected by the {\\it Fermi} satellite (GRBs 080916C, 090510, 090902B) and find that these photons were generated via synchrotron emission in the external forward shock. We arrive at this conclusion by four different methods as follows. (1) We check the light curve and spectral behavior of the $>$100 MeV data, and late time X-ray and optical data, and find them consistent with the so called closure relations for the external forward shock radiation. (2) We calculate the expected external forward shock synchrotron flux at 100 MeV, which is essentially a function of the total energy in the burst alone, and it matches the observed flux value. (3) We determine the external forward shock model parameters using the $>$100 MeV data (a very large phase space of parameters is allowed by the high energy data alone), and for each point in the allowed parameter space we calculate the expected X-ray and optical fluxes at late times (hours to days after the burst) and find these to be in good agreement with the observed data for the entire parameter space allowed by the $>$100 MeV data. (4) We calculate the external forward shock model parameters using only the late time X-ray, optical and radio data and from these estimate the expected flux at $>$100 MeV at the end of the sub-MeV burst (and at subsequent times) and find that to be entirely consistent with the high energy data obtained by {\\it Fermi}/LAT. The ability of a simple external forward shock, with two empirical parameters (total burst energy and energy in electrons) and two free parameters (circum-stellar density and energy in magnetic fields), to fit the entire data from the end of the burst (1--50s) to about a week, covering more than eight-decades in photon frequency --- $>$10$^2$MeV, X-ray, optical and radio --- provides compelling confirmation of the external forward shock synchrotron origin of the $>$100 MeV radiation from these {\\it Fermi} GRBs.  Moreover, the parameters determined in points (3) and (4) show that the magnetic field required in these GRBs is consistent with shock-compressed magnetic field in the circum-stellar medium with pre-shocked values of a few tens of micro-Gauss. ",
        "introduction": "The {\\it Fermi} Satellite has opened a new and sensitive window in the study of GRBs (gamma-ray bursts); for a general review of GRBs see Gehrels, Ramirez-Ruiz \\& Fox (2009), M\\'esz\\'aros (2006), Piran (2004), Woosley \\& Bloom (2006), Zhang (2007). So far, in about one year of operation, {\\it Fermi} has detected 12 GRBs with photons with energies $>$100 MeV. The $>$10$^2$ MeV emission of most bursts detected by the LAT (Large Area Telescope: energy coverage 20 MeV to $>$300 GeV) instrument aboard the {\\it Fermi} satellite shows two very interesting features (Omedei et al. 2009):  (1) The first $>$100 MeV photon arrives later than the first lower energy photon ($\\lae$1 MeV) detected by GBM (Gamma-ray Burst Monitor),  (2) The $>$100 MeV emission lasts for much longer time compared to the burst duration in the sub-MeV band (the light curve in sub-MeV band declines very rapidly).  \\begin{table*} \\begin{center} \\begin{tabular}{ccccccc} \\hline & $\\beta_{LAT}$ & $p$ & $z$ & $d_{L28}$ & $t_{GRB}[s]$ & $E_{\\gamma,iso}[erg]$  \\\\ \\hline \\hline GRB 080916C  & $1.20 \\pm 0.03$ & $2.4 \\pm 0.06$ & $4.3$ & $12.3$ & $60$ & $8.8\\times 10^{54}$\\\\ GRB 090510 & $1.1 \\pm 0.1$ & $2.2 \\pm 0.2$ & $0.9$ & $1.8$ & $0.3$ & $1.08\\times10^{53}$\\\\ GRB 090902B & $1.1 \\pm 0.1$ & $2.2 \\pm 0.2$ & $1.8$ & $4.3$ & $30$ & $3.63\\times10^{54}$\\\\ \\hline \\end{tabular} \\end{center} \\caption{{\\small The main quantities used in our analysis for these 3 GRBs. $\\beta_{LAT}$ is the spectral index for the $>100$MeV data, $p$ is the power-law index for the energy distribution of injected electrons i.e. $dn/d\\gamma\\propto \\gamma^{-p}$, $z$ is the redshift, $d_{L28}$ is the luminosity distance in units of $10^{28}$cm, $t_{GRB}$ is the approximate burst duration in the {\\it Fermi}/GBM band and $E_{\\gamma, iso}$ is the isotropic equivalent of energy observed in $\\gamma$-rays in the 10keV-10GeV band for GRB080916C and GRB090902B, and in the 10keV-30GeV band for GRB090510. Data taken from Abdo et al. (2009a, 2009b, 2009c), De Pasquale et al. (2010).}} \\end{table*}  There are many possible $>$100 MeV photons generation mechanisms proposed in the context of GRBs; see Gupta \\& Zhang (2007) and Fan \\& Piran (2008) for a review. Shortly after the observations of GRB 080916C (Abdo et al. 2009a), we proposed a simple idea: the $>$100 MeV photons in GRB 080916C are produced via synchrotron emission in the external forward shock (Kumar \\& Barniol Duran 2009). This proposal naturally explains the observed delay in the peak of the light curve for $>$100 MeV photons -- it corresponds to the deceleration time-scale of the relativistic ejecta -- and also the long lasting $>$100 MeV emission, which corresponds to the power-law decay nature of the external forward shock (ES) emission (the ES model was first proposed by Rees \\& M\\'esz\\'aros 1992, M\\'esz\\'aros \\& Rees 1993, Paczy\\'nski \\& Rhoads 1993; for a comprehensive review of the ES model, see, e.g., Piran, 2004, and references therein). Following our initial analysis on GRB 080916C, a number of groups have provided evidence for the external forward shock origin of {\\it Fermi}/LAT observations (Gao et al. 2009; Ghirlanda, Ghisellini, Nava 2010; Ghisellini, Ghirlanda, Nava 2010; De Pasquale et al. 2010).  In this paper we analyze the $>$100 MeV emission of GRB 090510 and GRB 090902B in detail, and discuss the main results of our prior calculation for GRB 080916C (Kumar \\& Barniol Duran 2009), to show that the high energy radiation for all these three arose in the external forward shock via the synchrotron process. These three bursts - one short and two long GRBs - are selected in this work because the high energy data for these bursts have been published by the {\\it Fermi} team as well as the fact that they have good afterglow follow up observations in the X-ray and optical bands (and also the radio band for GRB 090902B) to allow for a thorough analysis of data covering more than a factor 10$^8$ in frequency and $>10^4$ in time to piece together the high energy photon generation mechanism, and cross check this in multiple different ways.  In the next section (\\S2) we provide a simple analysis of the LAT spectrum and light curve for these three bursts to show that the data are consistent with the external forward shock model. This analysis consists of verifying whether the temporal decay index and the spectral index satisfy the relation expected for the ES emission (closure relation), and comparing the observed flux in the LAT band with the prediction of the ES model (according to this model the high energy flux is a function of blast wave energy, independent of the unknown circum-stellar medium density, and extremely weakly dependent on the energy fraction in magnetic fields).  We describe in \\S3 how the $>$100 MeV data alone can be used to theoretically estimate the emission at late times ($t\\gae$ a few hours) in the X-ray and optical bands within the framework of the external forward shock model, and that for these three bursts the expected flux according to the ES model is in agreement with the observed data in these bands.  Moreover, if we determine the ES parameters ($\\epsilon_e$, $\\epsilon_B$, $n$, and $E$)\\footnote{$\\epsilon_e$ and $\\epsilon_B$ are the energy fraction in electrons and magnetic field for the shocked fluid, $n$ is the number density of protons in the burst circum-stellar medium, and $E$ is the kinetic energy in the ES blast wave.} using only the late time X-ray and optical fluxes (and radio data), we can predict the flux at $>$100 MeV at any time after the deceleration time for the GRB relativistic outflow. We show in \\S3 that this predicted flux at $>10^2$MeV is consistent with the value observed by the {\\it Fermi} satellite for the bursts analyzed in this paper.  These exercises and results show that the high energy emission is due to the external shock as discussed in \\S3. We also describe in \\S3 that the magnetic field in the shocked fluid --- responsible for the generation of $>$ 100 MeV photons as well as the late time X-ray and optical photons via the synchrotron mechanism --- is consistent with the shock compression of a circum-stellar magnetic field of a few tens of micro-Gauss.  It is important to point out that we do not consider in this work the prompt sub-MeV emission mechanism for GRBs --- which is well known to have a separate and distinct origin as evidenced by the very rapid decay of sub-MeV flux observed by Swift and {\\it Fermi}/GBM (the flux in the sub-MeV band drops-off with time as $\\sim t^{-3}$ or faster as opposed to the $\\sim t^{-1}$ observed in the LAT band). Nor do we investigate the emission process for photons in the LAT band during the prompt burst phase. ",
        "conclusions": "The {\\it Fermi} Satellite has detected 12 GRBs with $>$100 MeV emission in about one year of operation.  In this paper we have analyzed the $>$100 MeV emission of three of them: two long-GRBs (090902B and 080916C) and one short burst (GRB 090510), and find that the data for all three bursts are consistent with synchrotron emission in the external forward shock.  This idea was initially proposed in our previous work on GRB 080916C (Kumar \\& Barniol Duran 2009), shortly after the publication of this burst's data by Abdo et al. (2009a). Now, there are three GRBs for which high energy data has been published, and for all of them we have presented here multiple lines of evidence that $>$100 MeV photons, subsequent to the prompt GRB phase, were generated in the external forward shock. The reason that high energy photons are detected from only a small fraction of GRBs observed by {\\it Fermi} is likely due to the fact that the high energy flux from the external forward shock has a strong dependence on the GRB jet Lorentz factor, and therefore very bright bursts with large Lorentz factors are the only ones detected by {\\it Fermi}/LAT (this was pointed out by Kumar \\& Barniol Duran 2009, who also suggested that there should be no difference in long and short bursts, as far as the $>$100 MeV emission is concerned - the high energy flux is only a function of burst energy and time, eq. 1).  We have analyzed the data in 4 different ways, and all of them lead to the same conclusion regarding the origin of $>$10$^2$MeV photons. First, we verified that the temporal decay index for the $>$100 MeV light curve and the spectral index are consistent with the closure relation expected for the synchrotron emission in the external forward shock. Second, we calculated the expected magnitude of the synchrotron flux at 100 MeV according to the external forward shock model and find that to be consistent with the observed value.  Third, using the $>$100 MeV data only, we determined the external shock parameters, and from these parameters we predict the X-ray and optical fluxes at late times and find that these predicted fluxes are consistent with the observed values within the uncertainty of our calculations, i.e. a factor of two (see figs. \\ref{090902B-fwd}, \\ref{080916C-fwd}). And lastly, using the late time X-ray, optical and radio fluxes --- which the GRB community has believed for a long time to be produced in the external forward shock --- we determine the external shock parameters, and using these parameters we predict the expected $>$100 MeV flux at early times and find the flux to be in agreement with the observed value (see fig. \\ref{090902B-reverse}). The fact that the $>$100 MeV emission and the lower energy ($\\lae 1$ MeV) emission are produced by two different sources at two different locations suggests that we should be cautious when using the highest observed photon energy and pair-production arguments to determine the Lorentz factor of the GRB jet.  We point out that the external shocks for these bursts were nearly adiabatic, i.e. radiative losses are small. The evidence for this comes from two different observations: (1) the late time X-ray spectrum lies in the adiabatic regime; (2) a radiative shock at early times (close to the deceleration time) would produce emission in the 10--10$^2$ keV band far in excess of the observed limits. We find that radiative shock is not needed to explain the temporal decay index of the $>$100 MeV light curve as suggested by (Ghisellini, Ghirlanda, Nava 2010), provided that the observing band is above all synchrotron characteristic frequencies.  We find that the magnetic field required in the external forward shock for the observed high and low energy emissions for these three bursts is consistent with shock-compressed magnetic field in the CSM; the magnetic field in the CSM -- before shock compression -- should be on the order of a few tens of micro-Gauss (see figs. \\ref{090902B-epsilonB}, \\ref{090902B-epsilonBlate} and \\ref{090510-epsilonB}). For these three bursts, at least, no magnetic dynamo is needed to operate behind the shock front to amplify the magnetic field.  The data for the short burst (GRB 090510) are consistent with the medium in the vicinity of the burst (within $\\sim$1 pc) being uniform and with density less than 0.1 cm$^{-3}$; the data rules out a CSM where $n\\propto R^{-2}$. On the other hand, the data for one of the two long {\\it Fermi} bursts (GRB 080916C) prefers a wind like medium and the other (GRB 090902B) a uniform density medium; these conclusions are reached independently from late time afterglow data alone and from the early time high energy data projected to late time using the 4-D parameter space technique described in \\S3.  It is also interesting to note that the power-law index of the energy distribution of injected electrons ($p$) in the shocked fluid, for all the three {\\it Fermi} bursts analyzed in this work, is $2.4$ to within the error of measurement, suggesting an agreement with the {\\it Fermi} acceleration of particles in highly relativistic shocks, e.g. Achterberg et al. (2001); a unique power-law index for electrons' distribution in highly relativistic shocks is not found in all simulations.  The study of high energy emission close to the deceleration time of GRB jets is likely to shed light on the onset of collisionless shocks and particle acceleration process.  It might seem surprising that we are able to fit all data (optical, X-ray, $\\lae10^2$MeV) for these three {\\it Fermi} bursts with just a few parameters for the external forward shock. This is in sharp contrast to Swift bursts which often display a variety of puzzling (poorly understood) features in their afterglow light curves. There are two reasons that these {\\it Fermi} bursts can be understood using a very simple model (external forward shock). (1) The data for the two long {\\it Fermi} bursts (080916C and 090902B) are not available during the first 1/2 day, and that is precisely the time frame when complicated features (plateau, etc., eg. Nousek et al. 2006, O'Brien et al. 2006) are seen in the X-ray afterglow light curves of Swift bursts (we note that the external forward shock model in its simplest form can't explain these features) --- however, the afterglow data at later times is almost invariably a smooth single (or double) power-law function that can be modeled by synchrotron emission from an external forward shock. (2) For very energetic GRBs --- the three bursts we have analyzed in this paper are among the brightest bursts in their class --- the progenitor star is likely to be completely destroyed leaving behind very little material to fall back onto the compact remnant at the center to fuel continued activity and give rise to complex features during the first few hours of the X-ray afterglow light curve (Kumar, Narayan \\& Johnson, 2008). To summarize, the GRB afterglow physics was simple in the decade preceding the launch of Swift, and then things became quite complicated, and now the {\\it Fermi} data might be helping to clear the fog and reveal the underlying simplicity once again."
    },
    "0910/0910.2555.txt": {
        "abstract": "IceCube is an all-flavor, cubic kilometer neutrino telescope currently under construction in the deep glacial ice at the South Pole.  Its embedded optical sensors detect Cherenkov light from charged particles produced in neutrino interactions in the ice.  For several years IceCube has been detecting muon tracks from charged-current muon neutrino interactions. However, IceCube has yet to observe the electromagnetic or hadronic particle showers or ``cascades'' initiated by charged-current or neutral-current neutrino interactions.  The first detection of such an event signature is expected to come from the known flux of atmospheric electron and muon neutrinos.  A search for atmospheric neutrino-induced cascades was performed using 275.46 days of data from IceCube's 22-string configuration.  Reconstruction and background rejection techniques were developed to reach, for the first time, a signal-to-background ratio $\\sim$1.  Above a reconstructed energy of 5~TeV, 12 candidate events were observed in the full dataset.  The signal expectation from the canonical Bartol atmospheric neutrino flux model is $5.63\\pm2.25$ events, while the expectation from the atmospheric neutrino flux as measured by IceCube's predecessor array AMANDA is $7.48\\pm1.50$ events.  Quoted errors include the uncertainty on the flux only.  While a conclusive detection can not yet be claimed because of a lack of background Monte Carlo statistics, the evidence that we are at the level of background suppression needed to see atmospheric neutrino-induced cascades is strong.  In addition, one extremely interesting candidate event of energy 133~TeV survives all cuts and shows an intriguing double pulse structure in its waveforms that may signal the ``double bang'' of a tau neutrino interaction.  \\abstractsignature ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.3111.txt": {
        "abstract": "% context heading (optional) {Methyl cyanide is an important trace molecule in star-forming regions. It is one of the more common molecules used to derive kinetic temperatures in such sources.} % aims heading (mandatory) {As preparatory work for Herschel, SOFIA, and in particular ALMA we want to improve the rest frequencies of the main as well as minor isotopologs of methyl cyanide.} % methods heading (mandatory) {The laboratory rotational spectrum of methyl cyanide in natural isotopic composition has been recorded up to 1.63~THz.} % results heading (mandatory) {Transitions with good signal-to-noise ratio could be identified for CH$_3$CN, $^{13}$CH$_3$CN, CH$_3^{13}$CN, CH$_3$C$^{15}$N, CH$_2$DCN, and $^{13}$CH$_3^{13}$CN in their ground vibrational states up to about 1.2~THz. The main isotopic species could be identified even in the highest frequency spectral recordings around 1.6~THz. The highest $J'$ quantum numbers included in the fit are 64 for $^{13}$CH$_3^{13}$CN and 89 for the main isotopic species. Greatly improved spectroscopic parameters have been obtained by fitting the present data together with previously reported transition frequencies.} % conclusions heading (optional) {The present data will be helpful to identify isotopologs of methyl cyanide in the higher frequency bands of instruments such as the recently launched Herschel satellite, the upcoming airplane mission SOFIA or the radio telescope array ALMA.} ",
        "introduction": "\\label{intro}  Methyl cyanide, CH$_3$CN, also known as acetonitrile, or cyanomethane, was among the early molecules to be detected by radio astronomical means. \\citet{det-MeCN} detected it almost 40 years ago toward the massive star-forming regions Sagittarius A and B close to the Galactic center. It has been detected in dark clouds such as TMC-1 \\citep{MeCN-TMC-1}, around the low-mass protostar IRAS 16293$-$2422 \\citep{MeCN_IRAS16293}, in the circumstellar shell of the famous carbon-rich star IRC~+10216 \\citep{MeCN_etc_IRC+10216}, in comets, such as Kohoutek \\citep{MeCN-Kohoutek}, and in external galaxies \\citep{MeCN_MeCCH_extragal}. As a strongly prolate symmetric top, meaning $A \\gg B$, transitions with different $K$ but the same $J$ occur in fairly narrow frequency regions, but sample rather different energies. Therefore, CH$_3$CN is quite commonly used to derive the kinetic temperatures of dense molecular clouds as already done in \\citet{det-MeCN} or in e.\\,g. \\citet{MeCN-T_kin}. In less dense clouds the large dipole moment of 3.922~D \\citep{MeCN-dipole} may prevent thermal equilibration. In that case, the isoelectronic methyl acetylene, or propyne, CH$_3$CCH, is often used instead.  The $^{12}$C/$^{13}$C ratio in the vicinity of the Galactic center is approximately 20 \\citep{isotopic_abundances,13C-VyCN_2008}. Thus, it is not surprising that the two isotopomers containing $^{13}$C were detected fairly soon after the detection of the main isotopolog. CH$_3 ^{13}$CN \\citep{MeCN-T_kin} was detected first probably by accident because the isotopic substitution at the central C-atom does not change the spectroscopic parameters much with respect to the main species. \\citet{Orion-survey_13CH3CN} detected lines of both isotopomers with one $^{13}$C in a line survey of Orion. More recently, CH$_2$DCN has been detected in two hot core sources \\citep{det-CH2DCN}, and some lines of CH$_3$C$^{15}$N were detected in a line survey of Sagittarius~B2 \\citep{Sgr-B2_Nummelin}.  A plethora of high-resolution spectroscopic investigations have been performed on methyl cyanide. It was actually among the first molecules to be studied by microwave spectroscopy by \\citet{MeCN_1st-MW}. \\citet{MeCN-LD} reviewed the investigations for the main isotopic species and performed high resolution, high accuracy measurements up to 1.2~THz. \\citet{MeCN_rot-isos_1996} analogously presented data up to 607~GHz for the singly substituted $^{13}$C isotopomers as well as for the species with $^{15}$N. \\citet{CH2DCN-rot} built on the very small data set for CH$_2$DCN by measuring an isotopically enriched sample up to 471~GHz. Finally, there has been only one report on the isotopolog with two $^{13}$C which was restricted to frequencies up to 72~GHz \\citep{MeCN-2x13}.  We have measured rotational spectra of methyl cyanide in natural isotopic composition both in wide frequency windows as well as selected lines up to 1.63~THz to provide new or updated catalog entries for various isotopologs of methyl cyanide as well as for excited vibrational states. With rotational temperatures of CH$_3$CN in hot cores reaching at least 150$-$200~K \\citep{Orion-survey_13CH3CN,det-PrCN_EtFo}, these data will be of great relevance for the Atacama Large Millimeter Array (ALMA). The main isotopolog and possibly even the $^{13}$C isotopomers or the $\\varv _8 = 1$ excited vibrational state will be seen with the HIFI instrument (Heterodyne Instrument for the Far-Infrared) on board of the recently launched Herschel satellite or with SOFIA (Stratospheric Observatory For Infrared Astronomy). In the present article we provide the ground state rotational data obtained for six different isotopic species: CH$_3$CN, $^{13}$CH$_3$CN, CH$_3^{13}$CN, CH$_3$C$^{15}$N, CH$_2$DCN, and $^{13}$CH$_3^{13}$CN; as usual, unlabeled atoms refer to $^{12}$C and $^{14}$N. Analyses of selected excited vibrational states are currently under way.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "\\label{conclusion} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  Rotational transitions for six astrophysically and astrochemically important isotopologs of methyl cyanide in their ground vibrational states have been recorded and the existing data sets have been extended greatly in most cases. The data should permit reliable predictions to above 2~THz in all instances, sufficient not only for Herschel or ALMA, but very likely also for such missions as SOFIA, CCAT, or others.  Transitions of the even rarer isotopomers containing $^{15}$N and one $^{13}$C may be observable also, but chance of significant overlap are even higher and the chances to observe these species in space seem vanishingly low.  Predictions of the rotational spectra of the isotopologs studied in the course of the present investigation will be available in the catalog section\\footnote{website: http://www.astro.uni-koeln.de/cdms/entries/, see also http://www.astro.uni-koeln.de/cdms/catalog/} of the Cologne Database for Molecular Spectroscopy\\footnote{website: http://www.astro.uni-koeln.de/cdms/} \\citep{CDMS_1,CDMS_2}. The complete line, parameter and fit files will be deposited in the Spectroscopy Data section of the CDMS. Updated or new JPL catalog \\citep{JPL-catalog} entries\\footnote{website: http://spec.jpl.nasa.gov/ftp/pub/catalog/catdir.html, see also http://spec.jpl.nasa.gov/} will be available also.   %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%   %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "0910/0910.1987_arXiv.txt": {
        "abstract": "Combination of Fresnel Zone Plates (FZP) can make an excellent telescope for imaging in X-rays. We present here the results of our experiments with several pairs of tungsten made Fresnel Zone plates in presence of an X-ray source kept at a distance of about $45$ feet. The quasi-parallel beam allowed us to study sources placed on the axis as well as off the axis of the telescope. We present theoretical study of the fringe patterns produced by the zone plates in presence of a quasi-parallel source. We compare the patterns obtained from experiments with those obtained by our Monte-Carlo simulations. The images are also reconstructed by deconvolution from both the patterns. We compare the performance of such a telescope with other X-ray imaging devices used in space-astronomy. ",
        "introduction": "\\label{intro}  Imaging in X-rays in space has always been a challenging task. X-ray films are generally very inefficient and in space it is inconvenient to take continuous imaging with X-ray films. To circumvent this, Coded Aperture Masks (CAM) (Mertz \\& Young, 1960; Ables, 1968; Dicke, 1968) are widely used. While a CAM has an advantage that it is a single element instrument, it has the disadvantage that the resolution is low, and is limited by the smallest mask element. Recently,  grazing incidence focusing instruments are being developed specially in space missions such as NUSTAR and SIMBOL-X (Ramsay, 2006; Pereschi, 2008), where target resolutions of $20-40$ arc seconds are being contemplated. A relatively simpler concept is, however, present in the literature. It was realized (Mertz 1965) that high resolution can be achieved when X-ray shadows are cast with two perfectly aligned Fresnel zone plates and the Moir\\'e fringes could be discerned by suitable detectors. Furthermore, the resolution is independent of the energy of X-rays. Theoretically, the resolution could be arbitrarily high as it increases with the separation $D$ of the zone plates, apart from being proportional to the finest element in the plate. Since small-pixel detectors were unavailable even though this concept was well established, there was hardly any application. Desai and his collaborators (1993, 1998, 2000) showed renewed interests in this type of telescopes and actually carried out an experiment and showed that an image reconstruction was possible.  However, so far, a detailed and quantitative exploration of imaging with single and multiple sources in laboratory environments has not been carried out with Fresnel Zone Plate telescopes. With the recent development of CMOS based X-ray detectors where each pixel could be as small as $50$ microns, it has become necessary to revisit the problem once more to actually study the feasibility of sending Zone Plate telescopes for space applications. Furthermore, the high resolution imaging devices could be used for medical purposes also. This is important since a high dose of X-ray is required when conventional X-ray films are used. In laboratory and medical uses, however, the beams are likely to be diverging. It is therefore necessary to study the feasibility of imaging with diverging beams as well.  In this series of papers, we plan to explore the practical aspects of the zone plate telescopes, focusing on the effects of the divergence of the beam and achievable resolutions in diverging beams. We also study with various combinations of zone plates. Our results are also supported by Monte-Carlo simulations carried out with parallel and divergent beams. Later in this series we will describe the results of the test and evaluation of two zone plate telescopes for the RT-2 experiments on board the Russian satellite CORONAS-PHOTON to be launched shortly.  The plan of our present paper is the following. In the next Section, we briefly describe the mathematical formalism behind the functioning of a Fresnel zone plate telescope and the properties of the Moir\\'e patterns falling on the detector. In Sections 3 and 4, we describe the setup of our experiment and the scheme of Monte-Carlo simulations. In Section 5, we present results obtained with a single on-axis or off-axis source with two different configurations of the zone plates and with two types of zone plates. Monte-Carlo simulations of these cases are also presented to support the experimental results. In Section 6, we present results with multiple sources and present simulation results. Finally, in Section 7, we make concluding remarks. In Paper II, of this series we will make a comparative study of Monte-Carlo simulations for single and multiple sources at various distances as well as for parallel beam, which is not achievable in the laboratories. In Paper III (Nandi et al., 2008), we will discuss the test and evaluations of the RT-2 payload to be launched aboard CORONAS-PHOTON satellite. ",
        "conclusions": "Zone plates have long been conceived to generate high resolution achromatic X-ray images. Its resolution is superior compared to that of the Coded Aperture Masks (CAMs) because of its finer zones. In the present paper, we concentrated on the results of an experimental set-up where the X-ray beam is almost parallel. The collimator is $45$ feet long and a zone plate of $15$ mm radius makes an angle of less then $4$ arc minutes at the source. We derived the nature of the fringes theoretically and showed that a point source will produce concentric circles on the detector plane. Through experiments and through Monte-Carlo simulations, we have been able to present the fringe patterns and also reconstructed the corresponding images both for the on-axis and off-axis sources. Partly because of the diverging nature of the beam and partly due to the finite size of the detector pixels, the ideal theoretical resolution was not achieved. This could however be improved during further image processing.  Resolution of a zone plate telescope could be made very high, since it is inversely proportional to the distance between the plates. From a practical point of view, in an astronomical satellite, a distance of a meter is possible which will yield (with smallest zones around $40-50$ micron as in our experiment) a ideal resolution of $8-10$ arc seconds. Another widely accepted type of X-ray imaging instrument having a very high resolution is the focusing grazing incidence type telescopes which are being contemplated in future space missions such as NUSTAR and SIMBOL-X. While these instruments target angular resolutions of $40$ and $20$ arc seconds respectively, to achieve these at high energies, formation flight architectures with focal lengths of $10$m and $20$m are required. Moreover, grazing incidence telescopes are expensive to build and the resolution is energy dependent, even when we assume the pointing is very accurate (any grazing type instrument is highly sensitive to alignment). In passing, we may remark that at a $20$m separation, our Zone plate telescope will definitely achieve an achromatic resolution of less than an arc second.  One disadvantage of the Zone plate telescopes is that if the size of the finest zone and the distance of separation of the plates are kept fixed, the resolution will not increase even if the areas of the plates are made bigger. However, a larger size will gather more photons and will increase the sensitivity of the instrument. An achievable zone plate size could be very big (say one meter diameter, made by low $Z$ metal quoted with high Z metal, such as lead or gold). Even with CZT detectors with 2.5mm pixel, reasonably high resolution would be achieved.  In order to check the efficiency of a zone plate telescope, we note that with our set up, simulations show (Paper-II) that a flux of about $10-12$ photons per cm$^2$ is required to have the source resolved about the background. With an increase in area this number does not change! For the crab nebula, the photon flux in $20-100$ keV range is about $0.12$ counts/sec/cm$^2$. Thus we need to integrate about $100$ seconds for imaging purpose. However, for a transient source, our instrument would be more effective since we could perhaps see change in size of the X-ray emitting region (such as glowing of the whole accretion disk when the accretion rate suddenly rises). With a larger field of view (as opposed to the focusing instruments), zone plate telescopes are ideally suited for imaging transient sources at a high energy. Of course, to eliminate high energy cosmic rays which will increase the background we require the usual protective shielding on our instrument.  In the next paper (Palit et al. 2008), of this series we shall concentrate on the image resolution, and how it depends on the source distance and the detector details. We shall also demonstrate the capability of a complete zone plate telescope which will constitute four sets of zone plates. In Paper III, we shall concentrate on the setups used in the RT-2/CZT-CMOS payload aboard the Russian solar satellite CORONAS-PHOTON (Nandi et al. 2008)."
    },
    "0910/0910.1349_arXiv.txt": {
        "abstract": "We present milliarcsecond-resolution radio very long baseline interferometry (VLBI) observations of the ultracool dwarfs TVLM\\,513--46546 (M8.5) and 2MASS J00361617+1821104 (L3.5) in an attempt to detect sub-stellar companions via direct imaging or reflex motion.  Both objects are known radio emitters with strong evidence for periodic emission on timescales of about 2 and 3 hours, respectively.  Using the inner seven VLBA antennas, we detect unresolved emission from TVLM\\,513--46546 on a scale of 2.5 mas ($\\sim 50$ stellar radii), leading to a direct limit on the radio emission brightness temperature of $T_B\\gtrsim 4\\times 10^5$ K.  However, with the higher spatial resolution afforded by the full VLBA we find that the source appears to be marginally and asymmetrically resolved at a low S/N ratio, possibly indicating that TVLM\\,513--46546 is a binary with a projected separation of $\\sim 1$ mas ($\\sim 20$ stellar radii).  Using the 7-hr baseline of our observation we find no astrometric shift in the position of TVLM\\,513--46546, with a $3\\sigma$ limit of about 0.6 mas. This is about 3 times larger than expected for an equal mass companion with a few-hour orbital period. Future monitoring of its position on a range of timescales will provide the required astrometric sensitivity to detect a planetary companion with a mass of $\\sim 10$ M$_J$ in a $\\gtrsim 15$ d ($\\gtrsim 0.06$ AU) orbit, or with a mass of $\\sim 2$ M$_J$ in an orbit of $\\gtrsim 0.5$ yr ($\\gtrsim 0.3$ AU). ",
        "introduction": "In recent years unexpectedly strong radio emission has been detected from very low mass stars and brown dwarfs (hereafter, ultracool dwarfs), revealing that these objects are capable of generating and dissipating kG-strength magnetic fields (e.g., \\citealt{ber01,ber06,hal08}).  The radio luminosity remains nearly uniform from early-M to mid-L dwarfs, even though other magnetic activity indicators (H$\\alpha$ and X-rays) decline by about two orders of magnitude over the same spectral type range (e.g., \\citealt{whw+04,ber06}).  Thus, radio observations provide a particularly powerful probe of the magnetic field properties.  The radio emission from several ultracool dwarfs also exhibits a periodicity that matches the stellar rotation velocity ($v\\,{\\rm sin} i$), thereby pointing to a large-scale, axisymmetric, and stable magnetic field configuration on timescales of hours to years \\citep{ber05,hal06,hal07,ber09}.  These conclusions are also supported by observations of Zeeman broadening in FeH molecular lines \\citep{rei07} and time-resolved optical spectropolarimetry \\citep{dmp08,mdp08}.  Since ultracool dwarfs are fully convective, the solar-type $\\alpha\\Omega$ dynamo, which operates at the shearing interface between the radiative and convective zones \\citep{par55}, cannot be responsible for generating and amplifying the inferred fields. However, recent numerical simulations suggest that large scale axisymmetric fields can still be generated in fully convective objects, at least for conditions that roughly correspond to mid-M dwarfs ($M\\sim 0.3$ M$_\\odot$; \\citealt{bro08}).  Simulations that correspond to the conditions in ultracool dwarfs are challenging and have not been investigated so far.  As a result, observational constraints on the scale, geometry, and origin of the fields are essential.  The spectroscopic Zeeman techniques are currently of limited utility for the dim and generally rapidly-rotating ultracool dwarfs, since they require very high signal-to-noise ratios and have reduced sensitivity at high rotation velocities due to line broadening \\citep{rei07}.  Detailed radio observations are thus crucial.  Here we present the first very long baseline interferometry (VLBI) observations of ultracool dwarfs, aimed at providing further constraints on their magnetic field properties.  In particular, these observations have sufficient angular resolution to detect astrometric shifts of a few tenths of a milliarcsecond that may be exerted by a sub-stellar companion with a period of a few hours.  They also allow us to directly image a radio-emitting companion down to milliarcsecond scales.  The presence of such putative companions may enhance the magnetic activity through direct or tidal interaction.  We detect the well-studied M8.5 dwarf TVLM\\,513-46546 with the VLBA, the first such detection for any ultracool dwarf.  The detectability of this object with the VLBA suggests that long-term astrometric monitoring could reveal the presence of a planetary-mass companion on a $\\gtrsim 15$ d orbit.  Such a detection would also usher in a new era of extrasolar planet studies through radio astrometry. ",
        "conclusions": "Using high angular resolution radio VLBI observations, we clearly detect the ultracool dwarf TVLM513--46 on an unprecedented small scale of about 2.5 mas or about 50 stellar radii.  The source is unresolved on this scale, but it may be marginally resolved on a $\\sim 1$ mas scale when using the longest VLBA baselines, although the S/N ratio is low.  Given the corresponding physical scale of about 20 stellar radii, this may point to a binary system with two radio-emitting sources in a $\\gtrsim 1$ d orbit.  Beyond the possibility of a spatially resolved binary, a search for an astrometric shift in the position of the source due to the gravitational influence of a close companion yields a $3\\sigma$ limit of $\\lesssim 0.6$ mas, about 3 times larger than the expected shift for an equal-mass companion in a few-hour orbit.  Since we lost the phased-VLA and GBT data, we expect that future VLBI observations will provide the requisite astrometric accuracy to search for an equal-mass companion in such a tight orbit.  These future observations will also allow us to determine whether the marginally-resolved emission from TVLM513--46 is indeed due to an equal mass companion with a projected separation of about 20 stellar radii.  If this is indeed the case, the expected maximum reflex motion is about 1 mas, easily detectable with VLBA+GBT observations that span several days.  Equally important, the fact that TVLM513--46 is detected with the VLBA opens the possibility to astrometrically search for a planetary-mass companion with a $\\gtrsim 10$ d orbit.  Monitoring of TVLM513--46 with the VLBA or VLBA+GBT on timescales of days to months to years will allow us to probe the existence of companions with masses of $\\sim 1-10$ M$_J$ and with orbits of $\\sim 15$ d to $\\sim 1$ yr.  These observations may lead to the first detection of an extrasolar planet via radio astrometry.  Thus, the use of VLBI observations opens up a wide companion parameter space for TVLM513--46, with detection via direct imaging for a roughly equal-mass radio-emitting companion down to a scale of $\\sim 20$ stellar radii, and detection via reflex motion on scales as small as a few stellar radii for an equal mass companion, and scales of $\\gtrsim 0.1$ AU for a planetary mass companion.  This first VLBA detection also paves the way for similar observations of known and future radio-emitting ultracool dwarfs in search of companions."
    },
    "0910/0910.2383_arXiv.txt": {
        "abstract": "Reported observations in H$\\alpha$, Ca II H and K or or other chromospheric lines of coronal rain trace back to the days of the Skylab mission. Offering a high contrast in intensity with respect to the background (either bright in emission if observed at the limb, or dark in absorption if observed on disk) these cool blobs are often observed falling down from high coronal heights above active regions. A physical explanation for this spectacular phenomenon has been put forward thanks to numerical simulations of loops with footpoint concentrated heating, a heating scenario in which cool condensations naturally form in the corona. This effect has been termed ``catastrophic cooling'' and is the predominant explanation for coronal rain. In this work we further investigate the link between this phenomenon and the heating mechanisms acting in the corona. We start by analyzing observations of coronal rain at the limb in the Ca II H line performed by the \\textit{Hinode} satellite. We then compare the observations with 1.5-dimensional MHD simulations of loops being heated by small-scale discrete events concentrated towards the footpoints (that could come, for instance, from magnetic reconnection events), and by Alfv\\'en waves generated at the photosphere. It is found that if a loop is heated predominantly from Alfv\\'en waves coronal rain is inhibited due to the characteristic uniform heating they produce. Hence coronal rain may not only point to the spatial distribution of the heating in coronal loops but also to the agent of the heating itself. We thus propose coronal rain as a marker for coronal heating mechanisms. ",
        "introduction": "Coronal loops are dynamical entities which exhibit heating and cooling processes constantly. Most of them are actually far from being in hydrostatic equilibrium \\citep{Aschwanden_2001ApJ...550.1036A}. The dynamical nature of loops can manifest in observations in a variety of forms, among which propagating intensity variations is a common one. Indeed, intensity variations traveling along coronal loops have been frequently observed in many different wavelengths, in hot coronal EUV lines as well as chromospheric cool lines. The agents producing these intensity variations can be either propagating waves or flows along coronal loops, and observers can have a hard time differentiating them despite their very different physical nature. Coronal rain is an example of intensity variations caused by flows. This spectacular phenomenon corresponds to cool plasma condensations (with chromospheric temperatures) in a hot environment (coronal temperatures) falling down along coronal loops, hence the very appropriate name. Coronal rain seems to be a common phenomenon above active regions, where loops are dense. Due to the low temperatures, the condensations appear as bright emission profiles in chromospheric lines such as H$\\alpha$ or Ca II H and K when observed at the limb \\citep[][Fig. \\ref{fig1}]{DeGroof_2004AA...415.1141D} or as dark absorption profiles when observed on disk. If observed in EUV wavelengths these propagating blobs appear as dark features \\citep{Schrijver_2001SoPh..198..325S}.  \\citet{DeGroof_2004AA...415.1141D, DeGroof05} report simultaneous observations in EIT 304 \\AA, H$\\alpha$ images from Big Bear Solar Observatory, and in 171 \\AA~ (\\textit{TRACE}) of intensity variations propagating along a coronal loop. Following a detailed analysis they put forward a series of points through which a differentiation between waves and flows (cool plasma condensations, coronal rain) can be done. Propagating waves frequently reported in EIT/\\textit{SOHO} 195 \\AA~ and \\textit{TRACE} 171 \\AA~ appear as low min-to-max intensity variations as compared to coronal rain. Most observed waves correspond to sound modes propagating at constant speed corresponding to the local sound speed of the corona, $\\sim$ 150 km s$^{-1}$. On the other hand, cool plasma condensations are seen to accelerate while falling along coronal loops to velocities above 100 km s$^{-1}$. Furthermore, waves have only been seen to propagate upwards, within the first 20 Mm of loops, after which they are usually damped, while cool condensations have only been seen to fall down along the loops from coronal heights. The falling speeds of the condensations have been reported to be lower than free fall speeds, due probably to the gas pressure along the loop, which increases considerably below the transition region. In some cases continuous flows from one footpoint to the other are also observed.  \\citet{Antiochos_1999ApJ...512..985A} propose a common physical mechanism for coronal rain and prominence formation. Ranges for temperatures and densities seem to be shared by both mechanisms, as well as the location of formation, high in the corona. While a prominence forms and is supported by the magnetic field for a time scale of days, coronal rain occurs on a time scale of minutes. It was shown that prominences may not need to rely on the geometry of the magnetic field to form (presence of ``dips''), contrary to the general belief. They showed that a coronal loop whose heating is concentrated towards the footpoints is subject to a thermal instability in the corona, which dramatically cools down to chromospheric temperatures in time scales of minutes once the density and temperatures have reached critical values. This phenomenon was termed ``catastrophic cooling'' and has so far gained acceptance as possible explanation for coronal rain. This mechanism occurs in a cyclic manner, matching reported periodicity of coronal rain observations \\citep{Schrijver_2001SoPh..198..325S}. This periodicity depends on geometrical aspects such as the loop length and the heating scale height, and can be obtained even with a time independent heating \\citep{Muller_2003AA...411..605M}. As the heating is concentrated towards the footpoints a larger mass flux is input into the corona as compared to uniform heating. The density of the corona increases in time while the energy flux is kept constant due to the constant heating input at the footpoints. This causes the loop to reach a maximum temperature after which it starts decreasing due to the decreasing of the heating per unit mass. The lower temperatures and higher densities increase the radiative losses which further cool the corona. Eventually, in a time scale of hours, the loop reaches a critical state and a thermal instability sets in. Then, in a time scale of minutes the temperature and density respectively decrease and increase dramatically to chromospheric values. The cool condensation subsequently falls down subject to gravity and plasma pressure. The loop subsequently is depleted and reheats rapidly due to the low density. The cycle then repeats. We will refer to these cycles as ``limit cycles\", as termed by \\citep{Muller_2003AA...411..605M}.  Observations seem to indicate that Coronal rain is a fairly common phenomenon of active regions, where loops are hot and dense. The coronal heating mechanism being a great unknown in solar physics, it is interesting to think about its underlying relationship to coronal rain. At a first glance coronal rain may appear as a random failure of the coronal heating mechanism to heat the loops in active regions. However, as previously stated, this phenomenon seems to act in the corona in a cyclic manner. Furthermore simulations have pointed out to the necessity of specific spatially dependent heating in order to allow the catastrophic cooling leading to coronal rain \\citep{Antiochos_1999ApJ...512..985A, Muller_2003AA...411..605M}. This cooling phenomenon seems then to be deeply linked to the unknown coronal heating mechanism. Since it is a fairly easily observable phenomenon due to the large velocities and density variation (hence clear Doppler shifted emission or absorption profiles), it may act as a marker of the operating heating mechanism in the loop. It has been shown for instance that footpoint concentrated heating may lead to catastrophic cooling if the heating scale height is sufficiently concentrated towards the footpoints \\citep{Muller_2003AA...411..605M}. Uniform heating on the other hand fails to reproduce the phenomenon since the heating rate per unit mass needs to decrease in time locally in the corona in order to allow the thermal instability to set in. It was further shown that catastrophic cooling does not need time dependent heating. In other words, it can happen even with a constant heating function. The footpoint concentrated heating function to which simulations point to matches the observational evidence of coronal loops above active regions for being mainly heated at their footpoints \\citep{Aschwanden_2001ApJ...560.1035A}, which sets most loops out of hydrostatic equilibrium. Further evidence of this fact has been found by \\citet{Hara_2008ApJ...678L..67H} using the Hinode/EIS instrument, which shows that active region loops exhibit upflow motions and enhanced nonthermal velocities. Possible unresolved high-speed upflows were also found, fitting in the footpoint concentrated heating scenario.  Many heating mechanisms have been proposed as candidates for heating the solar corona up to the observed few million degree temperatures. In this context a large emphasis has been put on the search for Alfv\\'en waves in the solar corona. Theoretically, they can be easily generated in the photosphere by the constant turbulent convective motions, which inputs large amounts of energy into the waves \\citep{Muller_1994AA...283..232M, Choudhuri_1993SoPh..143...49C}. Having magnetic tension as its restoring force the Alfv\\'en waves can travel less affected by the large transition region gradients with respect to other modes. Also, when traveling along thin magnetic flux tubes they are cut-off free since they are not coupled to gravity (Musielak et al. 2007)\\footnote{\\citet{Verth_etal_aap_09} have pointed out however that the assertion made by \\citet{Musielak_2007ApJ...659..650M} is valid only when the temperature in the flux tube does not differ from that of the external plasma. When this is not the case a cut-off frequency is introduced.}. Alfv\\'en waves generated in the photosphere are thus able to carry sufficient energy into the corona to compensate the losses due to radiation and conduction, and, if given a suitable dissipation mechanism, heat the plasma to the high million degree coronal temperatures \\citep{Uchida_1974SoPh...35..451U, Wentzel_1974SoPh...39..129W, Hollweg_1982SoPh...75...35H, Poedts_1989SoPh..123...83P, Ruderman_1997AA...320..305R, Kudoh_1999ApJ...514..493K, Antolin_2009arXiv0910.0962v1} and power the solar wind \\citep{Suzuki_2006JGRA..11106101S, Cranmer_2007ApJS..171..520C}. The main problem faced by Alfv\\'en wave heating is actually to find a suitable dissipation mechanism. Being an incompressible wave it must rely on a mechanism by which to convert the magnetic energy into heat. Several dissipation mechanisms have been proposed, such as parametric decay \\citep{Goldstein_1978ApJ...219..700G, Terasawa_1986JGR....91.4171T}, mode conversion \\citep{Hollweg_1982SoPh...75...35H, Kudoh_1999ApJ...514..493K, Moriyasu_2004ApJ...601L.107M}, phase mixing \\citep{Heyvaerts_1983AA...117..220H, Ofman_2002ApJ...576L.153O}, or resonant absorption \\citep{Ionson_1978ApJ...226..650I, Hollweg_1984ApJ...277..392H, Poedts_1989SoPh..123...83P, Erdelyi_1995AA...294..575E}. The main difficulty lies in dissipating sufficient amounts of energy in the correct time and space scales. For more discussion regarding this issue the reader can consult for instance \\citet{Klimchuk_2006SoPh..234...41K}, \\citet{Erdelyi_2007AN....328..726E} and \\citet{Aschwanden_2004psci.book.....A}. Works considering Alfv\\'en wave heating as coronal heating mechanism have shown that the obtained coronae are uniformly heated \\citep{Moriyasu_2004ApJ...601L.107M, Antolin_2008ApJ...688..669A, Antolin_2009arXiv0910.0962v1, Suzuki_2006JGRA..11106101S}. In these works the heating issues from shocks of longitudinal modes (mainly slow modes) from mode conversion of the Alfv\\'en waves due to the density fluctuation, wave-to-wave interaction and deformation of the wave shape during propagation. The coronal loops issuing from Alfv\\'en wave heating are found to satisfy quite well the RTV scaling law \\citep{Rosner_1978ApJ...220..643R} which quantifies the heating uniformity in the loops. This result would point then towards an inhibition of coronal rain if Alfv\\'en wave heating is predominant in the loop. It is this idea that is addressed in this work.  Another promising coronal heating candidate mechanism is the nanoflare reconnection heating model. The nanoflare reconnection process was first suggested by \\citet{Parker_1988ApJ...330..474P}, who considered a magnetic flux tube as being composed by a myriad of magnetic field lines braided into each other by continuous footpoint shuffling. Many current sheets in the magnetic flux tube would be created randomly along the tube that would lead to many magnetic reconnection events, releasing energy impulsively and sporadically in small quantities of the order of $10^{24}$ erg or less (nanoflares). Parker's original idea was a nanoflare reconnection heating acting uniformly in the corona but it was later proposed to be concentrated towards the footpoints of loops, where the magnetic canopy lies and magnetic field lines may entangle \\citep{Klimchuk_2006SoPh..234...41K, Aschwanden_2001ApJ...560.1035A}. In the reconnection scenario waves are also expected to be generated, and the energy imparted into Alfv\\'en waves is a matter of debate. The imparted energy may well depend on the location in the atmosphere of the reconnection event. \\citet{Parker_1991ApJ...372..719P} suggested a model in which 20~\\% of the energy released by reconnection events in the solar corona is transfered as a form of Alfv\\'en wave.\\citet{Yokoyama_1998ESASP.421..215Y} studied the problem simulating reconnection in the corona, and found that less than 10~\\% of the total released energy goes into Alfv\\'en waves. This result is similar to the 2-D simulation results of photospheric reconnection by \\citet{Takeuchi_2001ApJ...546L..73T}, in which it is shown that the energy flux carried by the slow magnetoacoustic waves is one order of magnitude higher that the energy flux carried by Alfv\\'en waves. On the other hand, recent simulations by Kigure et al. (private communication) show that the fraction of Alfv\\'en wave energy flux in the total released magnetic energy during reconnection in low $\\beta$ plasmas may be significant (more than 50~\\%). Since the observed ubiquitous intensity bursts (nanoflares) are thought to play an important role in the heating of the corona \\citep{Hudson_1991SoPh..133..357H} and since they are generally assumed to be a signature of magnetic reconnection it is then crucial to determine the energy going into the Alfv\\'en waves during the reconnection process. Moreover, \\citet{Moriyasu_2004ApJ...601L.107M} has shown that the observed spiky intensity profiles due to impulsive releases of energy may actually be a signature of Alfv\\'en waves. It was found that due to nonlinear effects Alfv\\'en waves can convert into slow and fast magnetoacoustic modes which then steepen into shocks and heat the plasma to coronal temperatures balancing losses due to thermal conduction and radiation. The shock heating due to the conversion of Alfv\\'en waves was found to be episodic, impulsive and uniformly distributed throughout the corona, producing an X-ray intensity profile that matches observations. Hence, \\citet{Moriyasu_2004ApJ...601L.107M} proposed that the observed nanoflares may not be directly related to reconnection but rather to Alfv\\'en waves.  Differentiating Alfv\\'en wave heating from nanoflare reconnection heating during observations is one of the main tasks needed in order to solve the coronal heating problem. Following the work of \\citet{Moriyasu_2004ApJ...601L.107M}, \\citet{Antolin_2008ApJ...688..669A} have compared both heating mechanisms by studying the hydrodynamic response of a loop subject to both kinds of heating. It was found that Alfv\\'en waves lead to a dynamic, uniformly heated corona with steep power law indexes (issuing from statistics of heating events) while nanoflare-reconnection heating leads to lower dynamics (besides the times when catastrophic cooling takes place in the case of footpoint concentrated heating) and shallow power laws. It was further found that footpoint nanoflare heating (namely, nanoflare reconnection heating concentrated towards the footpoints of loops) leads to hot upflows (as observed in the Fe\\,XV 284.16~\\AA\\ line) due to the plasma being heated rapidly towards the footpoints before being ejected into the corona, while Alfv\\'en wave heating leads to hot downflows due to the plasma achieving the maximum temperatures in the corona (rather than at the footpoints) and being carried back to the footpoints by the strong shocks generated there \\citet{Antolin_2009arXiv0903.1766A}. In this work we propose coronal rain as another observational signature through which both heating mechanisms can be distinguished.  We start first by reporting limb observations of coronal rain from \\textit{Hinode}/SOT in the Ca\\,II H line. The velocities and shapes of the falling condensations are analyzed. With a 1.5-dimensional code we then proceed to model a coronal loop being subject to a heating mechanism that is concentrated towards the footpoints, such as the nanoflare reconnection heating model \\citep[as proposed by][]{Aschwanden_2001ApJ...560.1035A, Klimchuk_2006SoPh..234...41K}. In the case considered catastrophic cooling happens 3 times and we select the first event to analyze and compare with our observations. Next we proceed by generating Alfv\\'en waves at the photosphere and gradually increase the amplitude. As the heating from the waves becomes non-negligible with respect to the heating flux from the nanoflare heating events catastrophic cooling is inhibited and the loop reaches a thermal equilibrium state. This result has important consequences since it points to the conclusion that Alfv\\'en waves are a non predominant heating mechanism in active region loops.  The work is organized as follows. In \\S\\ref{observations} we report observations of coronal rain in the Ca II H line performed with \\textit{Hinode}/SOT. In \\S\\ref{model} we introduce the 1.5-dimensional MHD model in which our loop is based on and discuss the heating models of the loop. In \\S\\ref{simulation} we present the results of footpoint heating and analyze a typical case of catastrophic cooling. The effect of Alfv\\'en waves on the thermal stability of the loop is also studied. In \\S\\ref{discussion} we discuss the results in the context of coronal heating and conclude the work. ",
        "conclusions": "\\label{discussion}  Observations in chromospheric lines such as H$\\alpha$ or Ca II H and K seem to show that coronal rain is a phenomenon exclusively of active regions \\citep{Schrijver_2001SoPh..198..325S, DeGroof_2004AA...415.1141D}, where loops are dense and heating appears to be concentrated towards the footpoints \\citep{Aschwanden_2001ApJ...550.1036A, Aschwanden_2001ApJ...560.1035A, Hara_2008ApJ...678L..67H}. In this work we have reported observations with the Hinode/SOT telescope in the Ca II H line of coronal rain occurring over an active region, and compared with results of simulations of loops that undergo catastrophic cooling. We have found that catastrophic cooling of coronal loops satisfactorily reproduces the main observed features of coronal rain. The loops are preferentially heated towards the footpoints and are subject to cycles \\citep[``limit cycles'', as termed by][]{Muller_2003AA...411..605M} in which they rapidly cool down and then reheat, they get dense and then deplete. The constant heating input at the footpoints of the loops produces coronae out of hydrostatic equilibrium. The coronal density increases in time causing a gradual decrease of the temperature and increase of the radiative losses. When the temperatures are sufficiently low radiative losses are dramatically increased since the condensation becomes optically thick and radiates considerably more. Catastrophic cooling sets in, either locally in the corona or globally, case in which the entire corona is cooled down to chromospheric values. Dense condensations of cool plasma form at coronal heights, which subsequently fall down by gravity. The loop is then depleted and the cycle starts again.  The characteristics of our condensation match well the main characteristics of the coronal rain from the reported observations. The velocities are in the same range as would be expected from coronal rain forming in roughly the same locations of loops having essentially the same length. We obtain an elongation of the condensation in the simulation as it gets denser and the effective gravity increases along its way down the loop. This elongation is observed as well in the present Hinode observations and has previously also been reported \\cite{Schrijver_2001SoPh..198..325S}. The temperatures and densities of the resulting condensation have chromospheric values, which suggest that they would emit radiation in lines such as H$\\alpha$ or Ca II H or K, as in the observations with Hinode/SOT or the Swedish Solar Telescope. In a future paper we will test this idea by synthesizing the H$\\alpha$ line profile based on self-consistent solution of the ionization rate equations coupled with the hydrodynamic equations.  As suggested by the simulations coronal rain is subject not only to gravity but also to the pressure changes inside the loop. Catastrophic cooling is the abrupt loss of thermal equilibrium in the corona, through which large pressure changes along the loop can occur and which can drive strong flows and shocks along the loop. In the catastrophic cooling event from the simulation reported here an upward motion of the condensation is obtained due to the local changes in pressure from the flows and the shocks. This may be a possible explanation to the upward motion of coronal rain as reported from the observations with Hinode/SOT in the Ca II H line.  If coronal rain is indeed the consequence of the catastrophic cooling mechanism we then must have a heating mechanism acting preferentially towards the footpoints in loops where coronal rain is observed, i.e. active region loops. We have seen that Alfv\\'en waves generated at the footpoints of loops produce uniform coronae and are thus unable to reproduce phenomena such as coronal rain. Furthermore, when Alfv\\'en waves are present in a loop and have enough energy flux to heat and maintain a hot corona the catastrophic cooling events are inhibited. The loop then converges to a uniform and steady state. Coronal loops in active regions show a recurrent occurrence of coronal rain. They are dynamical entities showing heating and cooling processes at all times. Our results indicate that having Alfv\\'en wave heating as the main heating mechanism in loops means the absence of coronal rain. These results indicate then that Alfv\\'en wave heating may not be the principal heating mechanism for coronal loops in active regions.  \\citet{Hara_2008ApJ...678L..67H}, using the \\textit{Hinode}/EIS instrument, have found upflow motions and enhanced nonthermal velocities in the hot lines of Fe XIV 274 and Fe XV 284 in active region loops. Possible unresolved high-speed upflows were also found. In \\citet{Antolin_2008ApJ...688..669A} and \\citet{Antolin_2009arXiv0903.1766A} we found that footpoint or uniform heating coming from nanoflare reconnection exhibit hot upflows, thus fitting in the observational scenario of active regions, while Alfv\\'en wave heating was found to exhibit hot downflows, which may fit in the observational scenario of quiet Sun regions \\citep{Chae_1998ApJS..114..151C, Brosius_2007ApJ...656L..41B}, further supporting the present conclusions. In \\citet{Antolin_2009arXiv0910.0962v1} we have found that Alfv\\'en wave heating is effective only in thick loops (with area expansions between photosphere and corona higher than 600) and long loops (with lengths of the corona above 50 Mm), a scenario which may not fit in active regions, where loops exhibit low area expansions due to the high magnetic field filling factors in those regions. Hence Alfv\\'en waves may play an important role in the heating of Quiet Sun regions (rather than active regions), where loops are often long, expand more than in active regions, kG (or higher) bright points are ubiquitous and coronal rain appears to be absent."
    },
    "0910/0910.1312_arXiv.txt": {
        "abstract": "We describe an extension of the hadronic SU(3) non-linear sigma model to include quarks. As a result, we obtain an effective model which interpolates between hadronic and quark degrees of freedom. The new parameters and the potential for the Polyakov loop (used as the order parameter for deconfinement) are calibrated in order to fit lattice QCD data and reproduce the QCD phase diagram. Finally, the equation of state provided by the model, combined with gravity through the inclusion of general relativity, is used to make predictions for neutron stars. \\vspace{1pc} ",
        "introduction": "Dense matter hadronic models, such as the ones used to describe neutron stars, work in a determined range of energies \\cite{Glendenning:1991ic,Weber:1989uq,Schaffner:1995th}. They can, in addition, fulfill some symmetries of the underlying theory (QCD) such as chiral symmetry \\cite{chiral2,Heide:1993yz,Carter:1995zi} and scale invariance \\cite{Bonanno:2008tt,Bonanno:2009fg} but they do not include deconfinement to quark matter. On the other hand, the broadly used bag-model \\cite{bag0} includes quark degrees of freedom but does not include chiral symmetry. Models such as the quark-NJL and quark sigma-models \\cite{Nambu:1961tp,Nambu:1961fr,Bub} include that symmetry, but do not directly incorporate hadronic degrees of freedom or quark confinement.  Our goal is to construct an effective model that contains hadronic and quark degrees of freedom present at different densities/temperatures but that can also coexist in a mixed phase. This allows the deconfinement phase transition to be a steep first order as well as a smooth crossover and cases in between. The last two possibilities, despite being predicted by lattice QCD, cannot be achieved by the usual procedure that connects, by hand, two different models with separate equations of state for hadronic and quark matter at the chemical potential at which the pressure of the quark EOS exceeds the hadronic one \\cite{hybrid1}.  In order to achieve this goal, we employ a single model for the hadronic and quark phases. Our extension of the  hadronic SU(3) non-linear sigma model also includes quark degrees of freedom in a spirit similar to the PNJL model \\cite{PNJL}, in the sense that it is a non-linear sigma model that uses the Polyakov loop $\\Phi$ as the order parameter for the deconfinement. This is a quite natural idea, since the Polyakov loop is related to the Z(3) symmetry, that is spontaneously broken by the presence of quarks. In QCD the Polyakov loop $\\Phi$ is defined via $\\Phi=\\frac13$Tr$[\\exp{(i\\int d\\tau A_4)}]$, where $A_4=iA_0$ is the temporal component of the SU(3) gauge field. ",
        "conclusions": "In spite of the fact that our model is relatively simple, it is the only one able to take into account different degrees of freedom and consequently allow steep as well as smooth transitions between different phases.  The model is in accordance with lattice QCD data and the phase diagram it reproduces is able to describe a broad variety of regimes, from compact stars to heavy ion collisions. Calculations along this line are in progress \\cite{soon}.  We conclude that our model is suitable for the description of neutron stars, since it predicts masses and radii in agreement with observations. Although it does not predict stable stars containing pure quark matter, the model allows stars that contain a core of mixed phase that can extend up to approximately $2$ km of radius. Even in this case, the reduced maximum star mass is still higher than the most massive pulsars observed.  A major advantage of our work compared to other studies of hybrid stars is that  we can study in detail the way in which chiral symmetry is restored and the way deconfinement occurs at high temperature/density. Since the properties of the physical system, e.g. the density of particles in each phase, are directly connected to the Polyakov loop it is not surprising that we obtain different results in a combined description of the degrees of freedom compared to a simple matching of two separate equations of state."
    },
    "0910/0910.3317_arXiv.txt": {
        "abstract": "\\noindent An analysis of the first set of low-redshift ($z$$<$$0.08$) Type~Ia supernovae monitored by the Carnegie Supernova Project between 2004 and 2006 is presented. The data consist of well-sampled, high-precision optical ($ugriBV$) and near-infrared (NIR; $YJHK_s$) light curves in a well-understood photometric system.  Methods are described for deriving light-curve parameters, and for building template light curves which are used to fit Type~Ia supernova data in the $ugriBVYJH$ bands.  The intrinsic colors at maximum light are calibrated using a subsample of supernovae assumed to have suffered little or no reddening, enabling color excesses to be estimated for the full sample. The optical--NIR color excesses allow the properties of the reddening law in the host galaxies to be studied.  A low average value of the total-to-selective absorption coefficient, $R_V \\approx 1.7$, is derived when using the entire sample of supernovae.  However, when the two highly reddened supernovae (SN~2005A and SN~2006X) in the sample are excluded, a value $R_V \\approx 3.2$ is obtained, similar to the standard value for the Galaxy. The red colors of these two events are well matched by a model where multiple scattering of photons by circumstellar dust steepens the effective extinction law.  The absolute peak magnitudes of the supernovae are studied in all bands using a two-parameter linear fit to the decline rates and the colors at maximum light, or alternatively, the color excesses.  In both cases, similar results are obtained with dispersions in absolute magnitude of 0.12--0.16~mag, depending on the specific filter-color combination. In contrast to the results obtained from the comparison of the color excesses, these fits of absolute magnitude give $R_V \\approx$ 1--2 when the dispersion is minimized, even when the two highly reddened supernovae are excluded.  This discrepancy suggests that, beyond the ``normal'' interstellar reddening produced in the host galaxies, there is an intrinsic dispersion in the colors of Type~Ia supernovae which is correlated with luminosity but independent of the decline rate.  Finally, a Hubble diagram for the best-observed subsample of supernovae is produced by combining the results of the fits of absolute magnitude versus decline rate and color excess for each filter.  The resulting scatter of 0.12~mag appears to be limited by the peculiar velocities of the host galaxies as evidenced by the strong correlation between the distance-modulus residuals observed in the individual filters.  The implication is that the actual precision of Type~Ia supernovae distances is 3--4\\%. ",
        "introduction": "\\label{sec:intro}  The discovery in the late-1990s from measurements of distant Type~Ia supernovae (SNe~Ia) that the expansion of the Universe is currently accelerating \\citep{riess98,perlmutter99} has given rise to an exciting new era in cosmology.  The observations appear to require a previously unrecognized form of energy density (``dark energy'') which today is the dominant constituent of the Universe. In recent years, observations of SNe~Ia to increasingly higher redshifts have confirmed this finding \\citep[e.g.,][]{riess07,astier06,wood-vasey07}. Measurements of the angular power spectrum of the cosmic microwave background radiation provide independent confirmation, favoring a flat geometry for the Universe with a dominant 73\\% of the energy density in the form of dark energy \\citep[e.g., see][]{spergel07}.  Baryon acoustic oscillations \\citep{eisenstein05} and measurements of the ratio of X-ray-emitting gas to total mass in galaxy clusters \\citep{allen04,allen07} have also provided convincing confirmation of this picture.  Effort is currently focused on determining the nature of dark energy by measuring an equation-of-state parameter of the form $w = P/(\\rho c^2$), and its time derivative $\\dot{w}$. Several ongoing and future projects aim at filling the Hubble diagram with high-redshift ($z>0.1$) SNe~Ia in order to measure $w$ and $\\dot{w}$ with sufficient precision to distinguish among the various proposed models of dark energy. However, these measurements require an improved reference set of SN~Ia observations at low redshift. Surprisingly, the currently available low-redshift datasets are neither sufficiently numerous nor homogeneous to adequately complement the high-redshift data.  The Carnegie Supernova Project (CSP) is one of several ongoing efforts to improve the quality of the low-redshift data.\\footnote[9]{See \\citet{contreras09} for a summary of other groups that are producing large databases of low-redshift SN~Ia light curves.}  Over the five-year period ending in May~2009, the CSP has obtained densely sampled optical {\\em and} near-infrared (NIR) light curves of $\\sim$100 SNe~Ia in a well-understood, homogeneous photometric system \\citep[][hereafter H06]{hamuy06}. The NIR was included specifically to address the unknown intrinsic colors and interstellar reddening to SNe~Ia.  This work presents an analysis of the first release of SNe~Ia by the CSP.  In \\S~\\ref{sec:lcs}, we present an analysis of the light curves, both optical and NIR, and introduce the methods used to derive parameters which serve to determine distances. In \\S~\\ref{sec:redd}, the non-trivial issue of measuring extinction is considered, and the nature of the reddening law in the host galaxies is studied. In \\S~\\ref{sec:absmag}, the properties of SNe~Ia as {\\em standardizable\\,} candles is examined by using Hubble-flow distances to fit the relationship between absolute peak magnitudes, decline rates, and colors or reddenings for different bands.  Finally, in \\S~\\ref{sec:concl}, the implications of our findings are discussed.  This paper is the second of three companion papers presenting first results from CSP observations of low- and high-redshift SNe~Ia.  The low-redshift data used in the present analysis are described in detail in the first paper by \\citet[][hereafter C09]{contreras09}.  In the third paper, NIR light curves of 35 high-redshift SNe~Ia observed with the Magellan Baade 6.5-m telescope are combined with the low-redshift data in order to construct the first rest-frame $i$-band Hubble diagram of SNe~Ia \\citep{freedman09}. ",
        "conclusions": "\\label{sec:concl}  This paper presents a first analysis of the light curves of 34 SNe~Ia followed by the CSP and released by C09. The high photometric precision and dense time sampling of the observations, especially in the optical bands, has allowed an in-depth examination of the general properties of the light curves of SNe~Ia in the $ugriBVYJH$ bands.  Subsamples of well-observed SNe --- 26 in the optical and 9 in the NIR ---  were used to build a family of template light curves using a technique that allows one to interpolate among the template data, taking into account variations in the sampling, in a 3-D space parameterized by the epoch relative to maximum light, the magnitude relative to maximum, and the decline rate.  The availability of observations covering a wide range of wavelengths (from $u$ to $K_s$) offers the opportunity to make further progress on the difficult issue of correcting SN luminosities for extinction outside the Galaxy.  However, depending on the approach taken, two quite different results are found.  In the first case, a subsample of 15 SNe assumed to have suffered little or no extinction in their host galaxies was used to measure color excesses for the whole sample.  By plotting the ratios of color excesses involving optical and NIR filters and comparing the values predicted by the CCM+O extinction law for interstellar dust, a value of the total-to-selective absorption coefficient of $R_V \\approx 1.7$ was derived, which is significantly lower than the Galactic average of $3.1$.  However, this value is largely influenced by two very red objects in the sample, SNe 2005A and 2006X.  When the same calculations are repeated after excluding these two objects, a value of $R_V=3.2\\pm0.4$ is found, in agreement with the Galactic average.  An alternative way of estimating $R_V$ is to express the absolute magnitudes of the SNe as a two-parameter function of the decline rate and color, and then to use the Hubble-flow distances (or Cepheid and SBF distances) of the SNe to derive the best fit parameters through $\\chi^2$ minimization. This method was first employed by \\citet{tripp98} and \\citet{tripp99}, who related the absolute magnitude in $B$ to the $B-V$ color and derived values of $R_V \\approx$ 1--1.5.  When this same technique is applied to the CSP sample of SNe~Ia using the peak magnitudes in $B$ and $J$ bands and the $B-V$ and $V-J$ colors, a value of $R_V \\sim$1--2 is obtained, regardless of whether the heavily reddened events, SNe 2005A and 2006X, are included or excluded. An equivalent analysis can also be carried out relating absolute magnitude to the decline rate and color excess.  Performing such fits for all available bands ($ugriBVYJHK_s$) again leads to systematically low values of the CCM+O-law parameter ($R_V \\approx 1.5$) regardless of whether the two highly reddened SNe are included in the sample.  These conflicting results on the value of $R_V$ reflect the variety of results found in the literature on this subject.  Most early attempts to derive the average properties of the host-galaxy reddening indicated that $R_V$ was unusually low \\citep[see][and references therein]{branch92}.  However, the available light curves at the time were largely photographic and there was considerable additional uncertainty regarding the nature of the intrinsic colors of SNe~Ia at maximum light.  The first modern study based on CCD data was made by \\citet{riess96}, who used the ratios of color excesses measured in $B-V$, $V-R$, and $V-I$ for a sample of 20 SNe~Ia to derive a value of $R_V = 2.55\\pm0.30$.  This analysis differed from previous studies in that the intrinsic color variation as a function of decline rate was taken fully into account in deriving the color excesses.  Similar values of $R_V$ have been found by \\citet{phillips99}, \\citet{altavilla04}, \\citet{reindl05}, and \\citet{wangx06}.  However, beginning with \\citet{tripp98}, studies based on the technique of minimizing the dispersion in the Hubble diagram of SNe~Ia have nearly uniformly led to values of $R_V$ in the range of $\\sim$1--2 \\citep[e.g.,][]{tripp99, astier06, conley07}.  If there is a pattern to these results, including our own in this paper, it would seem that attempts to measure $R_V$ via comparison of colors or color excesses tend to give larger values than the procedure of minimizing the scatter in the Hubble diagram with $R_V$ treated as a free parameter.  Nevertheless, there are exceptions to this rule such as the recent paper by \\citet{nobili08}, who studied the color evolution in $U-B$, $B-V$, $V-R$, and $R-I$ of 69 SNe~Ia with moderate reddening and obtained a value of $R_V = 1.01\\pm0.25$.  Common sense suggests that at least some portion of the observed reddening of SNe~Ia in spiral galaxies like the Milky Way must be due to dust in the interstellar medium of these galaxies.  Although only a few studies have been carried out of the nature of the reddening law in external galaxies, these are generally consistent with $R_V \\approx 3$, albeit with a large dispersion \\citep{goudfrooij94,patil07,ostman08}.  As shown in \\S~\\ref{sec:cered}, the ratios of the measured color excesses involving optical and NIR bands with respect to $E(B-V)$ are consistent with values of $R_V \\approx 3$.  As discussed by \\citet{krisciunas07}, inclusion of the NIR bands in this type of analysis improves considerably the accuracy with which $R_V$ can be measured.  Hence, it is tempting to conclude that we are measuring reddening produced by dust with very similar properties to that found in the interstellar medium of the disk of the Milky Way.  It will be interesting to see if this finding holds up as many more SNe~Ia with reliable optical$-$NIR color excesses are added to the CSP sample.  If we are, indeed, observing host-galaxy reddening due to ``normal'' Galactic-type dust, then the fact that low values of $R_V$ ($\\sim$1--2) are obtained when it is treated as a free parameter in minimizing the dispersion in the Hubble diagram is puzzling.  The implication is that there is an intrinsic dispersion in the colors of SNe~Ia which is correlated with luminosity, but is independent of the decline rate. Our finding in \\S~\\ref{sec:cered} that there is a small but systematic difference between color-excess measurements, $E(B-V)$, made at late epochs using the Lira law and those derived from the maximum light magnitudes, and that this difference correlates with color and perhaps also with absolute magnitude, may well be related to this.  Very red SNe potentially allow the dust reddening law to be studied on an individual basis.  Including the results for SN~2005A given in this paper, there are now six SNe~Ia for which a determination of $R_V$ has been made from optical and NIR photometry.  The results for these objects are summarized in Table~\\ref{tab:redsne}.  A weighted mean of the six determinations gives $R_V = 1.6$, with a surprisingly small rms of 0.3. The relatively narrow range of decline rates is also noteworthy, averaging $\\dm = 1.2$ with a dispersion of only 0.1.  Except for SN~2001el, the color excesses for these SNe are all very large.  \\citet{jha07} found that the distribution of $E(B-V)$ for SNe~Ia is well approximated by an exponential function with $\\tau = 0.138 \\pm 0.023$~mag.  Hence, values of $E(B-V) > 1$ mag should be extremely rare, although it must be kept in mind that the sample of SNe analyzed by Jha et al. was culled from several sources with different selection criteria.  Assuming that the Jha et al. distribution is correct, less than one SN for every thousand observed would be expected to have such a large amount of reddening --- yet according to the Asiago Supernova Catalog\\footnote[15]{http://web.oapd.inaf.it/supern/cat/ .}, the total number of nearby ($z < 0.02$) SNe~Ia discovered from 1985 through 2008 was only 358.  The fact that these heavily reddened SNe are much more common than expected, combined with the remarkable similarity of decline rates and $R_V$ values for the five objects in Table~\\ref{tab:redsne} with $E(B-V) > 1$ mag, leads us to speculate that these events actually represent a physically distinct subclass of SNe~Ia.  Hence, while extremely interesting in themselves, these objects like SNe 2005A and 2006X are not representative of the class of SNe~Ia employed in cosmological studies.  \\cite{wang05} has suggested that the presence of circumstellar dust distributed in a shell around the SN can produce reddening compatible with the observed low values of  $R_V$.  This hypothesis is supported by our observations of SNe~2005A and 2006X.  Specifically, we find that the red colors of these two SNe are better matched by  a model where multiple scattering of photons by circumstellar dust steepens the effective extinction law \\citep{goobar08}, than by the standard CCM+O Galactic reddening law with a low value of $R_V$.  Our results on the absolute peak magnitude calibrations for this new sample of CSP SNe confirm those obtained in previous studies, while extending the calibration to the optical $ugri$ and the NIR $YJHK_s$ bands.  Fits to the subsample of best-observed SNe employing the \\citet{tripp98} model, which assumes the absolute magnitudes to be a two-parameter function of the decline rate and color, give rms dispersions of 0.12--0.15~mag.  Remarkably, the same fits are found to apply equally well to either highly-reddened objects such as SNe 2005A and 2006X, or fast-declining, intrinsically-red events like 2005ke and 2006mr.  The simplicity of this method is appealing. We have also carried out fits which assume the absolute peak magnitudes to be a function of the decline rate and the measured color excesses.  These give quite similar results to those obtained with the Tripp model.  Fits to a large number of filters and colors give dispersions in the range of 0.12--0.20~mag for SNe with $0.7 < \\dm <1.7$ mag.  The quality of our NIR light curves allows, for the first time, detection of a weak dependence of the $J$-band luminosity on decline rate.  Combining the calibrations from all bands, we created a single Hubble diagram for the 23 best-observed SNe.  Although the resulting scatter of $0.12$~mag (6\\% in distance) is excellent, it does not represent a significant improvement over the precision obtained with individual filters.  This is attributed to the fact that, for any particular SN, there is a significant correlation in the error in the distance moduli obtained using different filters.  The source of these correlated errors appears to be the peculiar velocities of the SN host galaxies.  If true, this implies that the derived distances to the SNe are actually precise to 3--4\\%, making SNe~Ia competitive with even Cepheid variables as extragalactic distance indicators.  Finally, the set of NIR light curves presented here confirms the advantages of working in this wavelength region for cosmological studies.  Fit~8 of Table~\\ref{tab:tb} and Fit~13 of Table~\\ref{tab:ld} show that using either the $V-J$ color or color excess to correct the $J$ absolute magnitudes of the best-observed SNe (excluding the fast decliners) yields a dispersion of only 0.12~mag.  Moreover, this result is insensitive to the exact form of the reddening law since the extinction corrections in the NIR are small.  Note that the values derived for $\\sigma_{rm SN}$ of 0.02--0.04~mag suggest that the intrinsic dispersion in the $J$ band may be quite small.  This is obviously a preliminary result which must be treated with caution --- however, we will soon have a sample of $\\sim$80 SNe~Ia to test this. If confirmed, it implies that if the observational uncertainties can be reduced, there is much to be gained by extending the rest-wavelength coverage of future dark-energy experiments to include the $J$ band."
    },
    "0910/0910.3910_arXiv.txt": {
        "abstract": "We use our model for the formation and evolution of galaxies within a two-phase galaxy formation scenario, showing that the high-redshift domain typically supports the growth of spheroidal systems, whereas at low redshifts the predominant baryonic growth mechanism is quiescent and may therefore support the growth of a disc structure. Under this framework we investigate the evolving galaxy population by comparing key observations at both low and high-redshifts, finding generally good agreement. By analysing the evolutionary properties of this model, we are able to recreate several features of the evolving galaxy population with redshift, naturally reproducing number counts of massive star-forming galaxies at high redshifts, along with the galaxy scaling relations, star formation rate density and evolution of the stellar mass function. Building upon these encouraging agreements, we make model predictions that can be tested by future observations. In particular, we present the expected evolution to z=2 of the super-massive black hole mass function, and we show that the gas fraction in galaxies should decrease with increasing redshift in a mass, with more and more evolution going to higher and higher masses. Also, the characteristic transition mass from disc to bulge dominated system should decrease with increasing redshift. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.1915_arXiv.txt": {
        "abstract": "The single-degenerate channel is widely accepted as the progenitors of type Ia supernovae (SNe Ia). Following the work of Meng, Chen \\& Han (2009), we reproduced the birth rate and age of supernovae like SN 2006X by the single-degenerate model (WD + MS) with an optically thick wind, which may imply that the progenitor of SN 2006X is a WD + MS system. ",
        "introduction": "\\label{sect:1} Although type Ia supernovae (SNe Ia) showed their importance in determining cosmological parameters, e.g. $\\Omega_{\\rm M}$ and $\\Omega_{\\Lambda}$ ((Riess et al. 1998; Perlmutter et al. 1999), the progenitor systems of SNe Ia have not been confidently identified yet (Hillebrandt \\& Niemeyer 2000; Leibundgut 2000). Two basic scenarios for the progenitor of SN Ia have been discussed over the last three decades. One is a single degenerate (SD) model (Whelan \\& Iben 1973; Nomoto, Thielemann \\& Yokoi 1984), in which a CO WD increases its mass by accreting hydrogen- or helium-rich matter from its companion, and explodes when its mass approaches the Chandrasekhar mass limit. The companion may be a main-sequence star (WD + MS) or a red-giant star (WD + RG) (Yungelson et al. 1995; Li \\& van den Heuvel 1997; Hachisu et al. 1999a,b; Nomoto et al. 1999, 2003; Langer et al. 2000; Han \\& Podsiadlowski 2004; Chen \\& Li 2007, 2009; Han 2008; Meng, Chen \\& Han 2009; L\\\"{u} et al. 2009, Wang et al. 2009a, b). An alternative is the double degenerate (DD) model (Iben \\& Tutukov 1984; Webbink 1984), in which a system consisting of two CO WDs loses orbital angular momentum by gravitational wave radiation and merges finally. The merger may explode if the total mass of the system exceeds the Chandrasekhar mass limit (see the reviews by Hillebrandt \\& Niemeyer 2000 and Leibundgut 2000). In theory, a large amount of circumstellar materials (CSM) may form around SNe Ia via an optically thick wind for the SD model (Hachisu et al. 1996), while there is no CSM around DD systems. Then, a basic method to distinguish the two progenitor models is to find the CSM around progenitor systems. The CSM may play a key role to solve the problem of the low value of reddening ratio (Wang 2005), which is very important for precision cosmology (Wang et al. 2008). In addition, it is possible for the CSM to be the origin of color excess of SNe Ia (Meng et al. 2009).  Evidence for CSM was first found in SN2002ic (Hamuy et al. 2003), which has shown extremely pronounced hydrogen emission lines which have been interpreted as a sign of strong interaction between supernova ejecta and CSM. The discovery of SN2002ic may uphold the SD model (Han \\& Podsiadlowski 2006). The evidence for CSM was also found in a normal SN Ia (SN 2006X) defined by Branch, Fisher \\& Nugent (1993) and the CSM is proposed to be from a wind from a red-giant companion (Patat et al. 2007), while Hachisu et al. (2008) argued a WD + MS nature for this SN Ia. Blondin et al. (2009) found a similar signal to that of SN 2006X in another SNe Ia (SN 1999cl) and their results indicated that the supernovae like SN 2006X are rare objects ($2/31\\sim6\\%$). Recently, the third example like SN 2006X (SN 2007le) was reported (Simon et al. 2009), and then the ratio of 2006X-like supernova is increased to $3/32\\sim9\\%$.  In theory, SNe Ia can explode during the optically thick wind phase or after the optically thick wind while in stable or unstable hydrogen-burning phase (Han \\& Podsiadlowski 2004). The materials lost as the optically thick wind form CSM (Hachisu et al. 2008; Meng et al. 2009). If a SN Ia explodes after the optically thick wind, the CSM have been dispersed too thin to be detected immediately after the SN Ia explosion (Hachisu et al. 2008; Meng, Yang \\& Geng 2009). If a SNe Ia explodes during the optically thick wind phase, the materials lost from system form CSM very near the SN Ia, which may show the signal similar to SN 2006X (Hachisu et al. 2008; Meng, Yang \\& Geng 2009). Recently, Meng, Chen \\& Han (2009) performed binary stellar evolution calculations for more than 25,000 close WD + MS binary systems with various metallicities. In their works, the prescription of Hachisu et al. (1999a) for the accretion efficiency of hydrogen-rich material was adopted by assuming an optically thick wind (Hachisu et al. 1996), and then their works provide a possibility to check whether the SD model with optically thick wind can explain the birth rate of supernovae like SN 2006X. The purpose of this paper is to check the possibility by a binary population synthesis (BPS) approach.  In section \\ref{sect:2}, we simply describe our method. We show the results in section \\ref{sect:3} and give discussions and conclusions in section \\ref{sect:4}. ",
        "conclusions": "\\label{sect:4} \\subsection{the age of SN 2006X} \\label{sect:4.1}  SN 2006X is the first case that show a variable Na I D line in its spectral, which clearly shows a signal of CSM (Patat et al. 2007). Patat et al. (2007) also noticed a relatively low expansion velocity of the CSM ($\\sim50$ ${\\rm km/s}$). The low expansion velocity seems to imply that the progenitor of SN 2006X belongs to a WD + RG systems. However, Hachisu et al. (2008) showed that under the assumption of the optically thick wind, the WD + MS system also may produce the low-velocity CSM. In this paper, we show that if SN 2006X is from a WD + MS system with the optically thick wind, the age of its progenitor is smaller than 1 Gyr, which implies that there is star formation during the recent 1 Gyr in its host galaxy. The host galaxy of SN 2006X, NGC 4321 (M100), is a well-studied spiral galaxy (e.g. Ho et al. 1997) and SN 2006X locates near one of its arms (Wang et al. 2008). As one of the largest spiral galaxies in the Virgo cluster, it produced SNe 1901B, 1914A, 1959E, 1979C, and 2006X in roughly the last century and shows significant signal of star formation at present (Kanpen et al. 1993, 1996). Recently, Blondin et al. (2009) found a similar signal to that of SN 2006X in another SNe Ia (SN 1999cl) in archives. Its host galaxy, NGC4501 (M88), is also a spiral galaxy and SN 1999cl locates at one of its arms (Krisciunas et al. 2000). During the last 1 Gyr, the star formation is significant in the galaxy (Wong \\& Blitz 2002). The host galaxy of SN 2007le (NGC 7721) is a Sc galaxy, and a recent star formation is also expected (Iglesias-P\\'{a}ramo et al. 2006; Simon et al. 2009). So, all the supernovae fulfill the age constraint from WD + MS system. \\subsection{Progenitor system} \\label{sect:4.2}  Following the study of Meng, Chen \\& Han (2009), in this paper, we show the evolution of birth rate of supernovae like SN 2006X by assuming that if a SN Ia explodes at the optically thick wind phase, it is a 2006X-like supernova. Based on the assumption, we can reproduce the birth rate of 2006X-like objects. Please keep in mind that we only consider the case of WD + MS channel in this paper. However, the progenitor of SN 2006X is still an open question. Patat et al. (2007) suggested that the progenitor of SN 2006X belongs to WD + RG system such a RS Oph based on the velocity of absorptions line of Na I D. L\\\"{u} et al. (2009) designed a WD + RG model with an aspherical stellar wind and equatorial disk, and then they may explain some properties of SN 2006X. However, the birth rate of 2006X-like objects obtained in L\\\"{u} et al. (2009) is too low (less than 1\\%) to compare with that observed. Based on the assumption of the optically thick wind and a mass-stripping effect, Hachisu et al. (2008) argued that the progenitor of SN 2006X should be a WD + MS system and they also can well explain the properties of SN 2006X. Although the treatment of binary evolution in Meng, Chen \\& Han (2009) is different from that in Hachisu et al. (2008), Meng, Chen \\& Han (2009) obtained similar results for the SNe Ia exploding at the optically thick wind phase, e.g. there is no supernova explosion at the optically thick wind phase when the initial mass of CO WD is smaller than 0.9 $M_{\\odot}$. So, considering the results in this paper, all the properties of SN 2006X can be explained by the WD + MS channel with the optically thick wind, including its birth rate and its age. Then, we suggest that the progenitor of supernovae like SN 2006X is WD + MS system, not WD + RG system."
    },
    "0910/0910.4192_arXiv.txt": {
        "abstract": "The Fermi Gamma-ray Space Telescope (FGST) has opened a new high-energy window in the study of Gamma-Ray Bursts (GRBs).  Here we present a thorough analysis of GRB~080825C, which triggered the Fermi Gamma-ray Burst Monitor (GBM), and was the first firm detection of a GRB by the Fermi Large Area Telescope (LAT).  We discuss the LAT event selections, background estimation, significance calculations, and localization for Fermi GRBs in general and GRB~080825C in particular. We show the results of temporal and time-resolved spectral analysis of the GBM and LAT data. We also present some theoretical interpretation of GRB~080825C observations as well as some common features observed in other LAT GRBs. ",
        "introduction": "Gamma-Ray Bursts (GRBs) originate from the most luminous explosions in the universe and more than 35 years after their discovery in 1967 \\citep{kel73}, many questions remain to be answered about their possible progenitors, the composition of the ultra-relativistic outflows that power them, and the dominant emission mechanism for their prompt gamma rays.  The Burst And Transient Source Experiment (BATSE) onboard the Compton Gamma-Ray Observatory (CGRO; 1991-2000) made significant advances to the field, thoroughly exploring the 25~keV -- 2~MeV energy range with detailed population studies of the prompt gamma-ray emission.  Burst spectra were found to be well described by the Band function \\citep{band93}, which consists of two smoothly connected power laws. It was understood, however, that observations of GRBs at higher energies were of crucial importance to resolve some of the open issues: constrain the bulk Lorentz factor of the outflow and the distance from the central source to the gamma-ray emission region, distinguish between hadronic and leptonic origins of the gamma-ray emission, and probe for signatures of Ultra High Energy Cosmic Rays (UHECRs) which could be accelerated within GRB jets \\citep[see][for a review of the prospects for GRB science with Fermi LAT]{band08}.  Constraints on the origin of the high-energy emission from GRBs are quite limited due to both the small number of bursts with firm high-energy detection and the small number of events that were detected in such cases.  High-energy emission from GRBs was first observed by the Energetic Gamma-Ray Experiment Telescope (EGRET, covering the energy range from $30\\;$MeV to $30\\;$GeV) onboard CGRO. Emission above $100\\;$MeV was detected in five cases: GRBs~910503, 910601, 930131, 940217 and 940301 \\citep{dingus95}. One of these sources, GRB~930131, had high-energy emission that was consistent with an extrapolation from its spectrum obtained with BATSE between 25~keV -- 4~MeV~\\citep{som94}.  In contrast, evidence for an additional high-energy component up to $200\\;$MeV with a different temporal behavior to the low-energy component was discovered in GRB~941017 \\citep[in EGRET's calorimeter TASC; ][]{gon03}. The high-energy emission for the latter GRB lasted more than 200 seconds with a single spectral component being ruled out. A unique aspect of the high-energy emission in GRB~940217 was its duration, which lasted up to $\\sim$$90$ minutes after the BATSE GRB trigger, including an 18 GeV photon at $\\sim$$75$ minutes post-trigger \\citep{hur94}. More recently, the GRID instrument onboard Astro-rivelatore Gamma a Immagini LEggero (AGILE) detected 10 high-energy events with energies up to $300\\;$MeV from GRB 080514B, in coincidence with its lower energy emission, with a significance of $3.0\\;\\sigma$ \\citep{giu08}.  The Fermi Gamma-ray Space Telescope was launched on June 11 2008 and provides an unprecedented energy coverage and sensitivity for the study of high-energy emission in GRBs. It is composed of two instruments: the Gamma-ray Burst Monitor \\citep[GBM;][]{meg08} and the Large Area Telescope \\citep[LAT;][]{atw09}. The GBM covers the entire unocculted sky with 12 sodium iodide (NaI) detectors with different orientation placed around the spacecraft and covering an energy range from 8 keV to 1 MeV, and two bismuth germanate (BGO) scintillators placed on opposite sides of the spacecraft with energy coverage from 200 keV to 40 MeV. The LAT is a pair conversion telescope made up of $4 \\times 4$ arrays of silicon strip trackers and cesium iodide calorimeter modules covered by a segmented anti-coincidence detector designed to efficiently reject charged particles. The energy coverage of the LAT instrument ranges from 20 MeV to more than 300~GeV with a field-of-view (FoV) of $\\sim$$2.4$ steradians. Note that the LAT effective area is still non-zero even as far out as 70 degrees off-axis which allows the detection of bursts with such high incident angles. As of June 1 2009, 9 GRBs have been detected by the LAT at energies above 100~MeV: GRB~080825C \\citep{bou08}, GRB~080916C \\citep{abd09, tajima08}, GRB~081024B \\citep{omo08}, GRB~081215A \\citep{mcenery08}, GRB~090217 \\citep{ohn09}, GRB~090323 \\citep{ohnovdh09}, GRB~090328 \\citep{mcenery09, cutini09}, GRB~090510 \\citep{ohnopela09, omo09}. In this paper, we report the observations and analyses of gamma-ray emission from GRB~080825C, the first GRB detected by both the GBM and the LAT instruments. Section \\ref{sec:observations} will present the GBM and LAT observations along with the various methods used for data analysis, section \\ref{sec:dataana} provides the results of detailed time-resolved spectroscopy, and section \\ref{sec:discussion} discusses the theoretical interpretation of our observations and compares the properties of this event to the ones observed in some other LAT GRBs. ",
        "conclusions": "GRB~080825C is the first GRB detected by the LAT, with 13 events above 80~MeV and a detection significance of $\\sim$$6\\;\\sigma$.  The highest energy events, up to $\\sim$$600\\;$MeV, are detected at late times, $\\sim$$25-35\\;$s after the GBM trigger, when the emission in GBM has decreased close to the background level.  The lack of $>1$~GeV events in the LAT and the large angle of the source to the LAT boresight result in a localization uncertainty of $1.1^\\circ$ (statistical plus systematic) at the $1\\;\\sigma$ level.  The prompt emission spectrum from both instruments onboard Fermi covers over 5 decades in energy.  We have performed time-resolved spectral analysis using the two GBM NaI detectors with the brightest GRB signal, both GBM BGO detectors, and for the LAT with an event selection scheme that is optimized for GRB analysis.  We have carefully taken the energy-dependent backgrounds into account for both GBM and LAT, and studied the systematic uncertainties in the spectral analysis. The time-integrated and time-resolved spectra are well fit by the Band function with a hard-to-soft evolution in the first $25\\;$s: $E_{\\rm{peak}}$ evolves from $\\sim$$300$ to $\\sim$$150\\;$keV, the high-energy power-law index $\\beta$ is constant at a value of $\\sim$$-2.5$, while the low-energy power-law index $\\alpha$ is fairly constant except for the second time bin which contains the first LAT events.  In the last time bin, $\\sim$$25-35\\;$s after the GBM trigger time, the GBM data are barely above background level, and the spectrum is best fit by a single power law with an index of $\\sim$$-2$ which is significantly harder than the $\\beta$ values of the earlier intervals.  The duration and start time of the late broad peak in the high-energy emission, $\\sim$$16-31\\;$s after the trigger, suggest that this peak is emitted by the external reverse or forward shocks, rather than by internal dissipation within the GRB outflow (e.g., internal shocks or magnetic reconnection). The relatively fast flux decay after this peak slightly favors a reverse-forward shock `external' Compton origin over a forward shock SSC origin. Although the origin and emission mechanism for this late peak cannot be conclusively determined because of low number statistics (and the lack of observations at X-ray or optical wavelengths, due to the poor GRB localization), the external shock origin is further supported by the change in spectral behavior, in particular of the spectral index, at these late times. Observations of more, brighter GRBs with both GBM and LAT will be able to test this hypothesis."
    },
    "0910/0910.4471_arXiv.txt": {
        "abstract": "{% Nova V5116 Sgr 2005 No. 2, discovered on 2005 July 4, was observed with XMM-Newton in March 2007, 20 months after the optical outburst. The X-ray spectrum showed that the nova had evolved to a pure supersoft X-ray source, indicative of residual H-burning on top of the white dwarf. The X-ray light-curve shows abrupt decreases and increases of the flux by a factor 8 with a periodicity of 2.97~h, consistent with the possible orbital period of the system. The EPIC spectra are well fit with an ONe white dwarf atmosphere model, with the same temperature both in the low and the high flux periods. This rules out an intrinsic variation of the X-ray source as the origin of the flux changes, and points to a possible partial eclipse as the origin of the variable light curve. The RGS high resolution spectra support this scenario showing a number of emission features in the low flux state, which either disappear or change into absorption feature in the high flux state. A new XMM-Newton observation in March 2009 shows the SSS had turned off and V5116 Sgr had evolved into a weaker and harder X-ray source.} ",
        "introduction": "V5116 Sgr was discovered as Nova Sgr 2005b on 2005 July 4.049~UT, with magnitude $\\sim$8.0, rising to mag 7.2 on July 5.085 \\cite{lil05}. It was a fast nova ($t_3=39$ days) belonging to the Fe II class in the Williams (1992) classification (Williams et al. 2008). The expansion velocity derived from a sharp P-Cyg profile detected in a spectrum taken on July 5.099 was $\\sim$1300~km/s. IR spectroscopy on July 15 showed emission lines with FWHM~$\\sim$2200~km/s \\cite{rus05}. Photometric observations obtained during 13 nights in the period August-October 2006 show the optical light curve modulated with a period of $2.9712\\pm0.0024$~h \\cite{dob07}, which the authors interpret as the orbital period. They propose that the light-curve indicates a high inclination system with an irradiation effect on the secondary star. The estimated distance to V5116~Sgr from the optical light curve is $11\\pm3$~kpc \\cite{sal08}. A first X-ray observation with Swift/XRT (0.3--10~keV) in August 2005 yielded a marginal detection with 1.2($\\pm$1.0) $\\times10^{-3}$~cts/s \\cite{nes07a}. Two years later, on 2007 August 7, Swift/XRT showed the nova as a SSS, with 0.56 $\\pm$0.1~cts/s \\cite{nes07b}. A first fit with a blackbody indicated $T\\sim4.5\\times10^5K$. A 35~ks Chandra spectrum obtained on 2007 August 28 was fit with a WD atmospheric model with N$_{\\rm H}=4.3\\times10^{21}$~cm$^{-2}$ and $T=4.65\\times10^5K$ \\cite{NelOri07}. ",
        "conclusions": ""
    },
    "0910/0910.3856_arXiv.txt": {
        "abstract": "We discuss three effects of axial rotation at low metallicity. The first one is the mixing of the chemical species which is predicted to be more efficient in low metallicity environments. A consequence is the production of important quantities of primary $^{14}$N, $^{13}$C, $^{22}$Ne and a strong impact on the  nucleosynthesis of the  {\\it s}-process elements. The second effect is a consequence of the first. Strong mixing makes possible the apparition at the surface of important quantities of CNO elements. This increases the opacity of the outer layers and may trigger important mass loss by line driven winds. The third effect is the fact that, during the main-sequence phase, stars, at very low metallicity, reach more easily than their more metal rich counterparts, the critical velocity\\footnote{The critical velocity is the surface equatorial velocity such that the centrifugal acceleration compensates for the local gravity.}. We discuss the respective importance of these three effects as a function of the metallicity. We show the consequences for the early chemical evolution of the galactic halo and for explaining the CEMP stars. We conclude that rotation is probably a key feature which contributes in an important way to shape the evolution of the first stellar generations in the Universe. ",
        "introduction": "The study of the most iron poor stars in the halo offers a unique window on the nature and the evolution of the first generations of stars which provided the matter from which these halo stars were formed. Constraints on the first stellar generations may provide interesting views on topical questions as the nature of the stars which reionize the Universe, the initial mass function of the first stellar generations, the frequency of long soft Gamma Ray Bursts at very low metallicity, the timescale for the mixing of the newly synthesized products with interstellar medium material,  or the conditions of star formation in the very early life of our Galaxy.  \\begin{figure}[t] \\includegraphics[width=2.4in,height=2.4in,angle=0]{chem5-2.eps} \\hfill \\includegraphics[width=2.4in,height=2.4in,angle=0]{chem5N-1.eps} \\caption{{\\it Left panel} : variation of the abundances of various elements  (in mass fraction) as a function of the Lagrangian mass in the outer layers of a 60 M$_\\odot$ model at $Z=10^{-5}$ with $\\upsilon_{\\rm ini}=800$ km s$^{-1}$ at the end of the core He-burning phase. The grey area covers the CO core. {\\it Right panel :} same as left part for a non rotating 60 M$_\\odot$ model at $Z=10^{-5}$. Models computed by \\citet{paperVIII2002} and \\citet{Meynet2006}.} \\label{struc} \\end{figure}  In this context, many observed features of the very iron poor stars are quite surprising and were not at all expected: \\begin{itemize} \\item One expected that halo stars, having a very low [Fe/H], should show an important scatter in their abundances. Indeed, at these very early times, stars are expected to form from not well mixed clumps and thus should bear the nucleosynthetic signatures of a few peculiar events. But \\citet{Cayrel2004} find that the scatter in the abundances of several elements ratios (e.g. [Cr/Fe]) is very small. It can be as low as 0.05 dex. In contrast, r-process elements and nitrogen for instance reveal large star-to-star scatters \\citep{Ryan1996, Honda2004, Andrievsky2009}. Explanation are proposed in  \\citet{Ishimaru2004, Ishimaru2006, Cescutti2008}. \\item One expected that the very first generations of Pop. III stars were very massive and contributed to the enrichment of the interstellar medium through Pair Instability Supernovae \\citep[PISN,][]{Barkat1967, Bond1984}. The nucleosynthetic pattern produced by such events is well understood and presents specific features  \\citep{HegerWoosley2005}, which were expected to be observed at least in some iron-poor halo stars. At the moment no trace of PISN has been found  \\citep{Cayrel2004}.  Some authors have suggested an explanation why this signature has, up to now, escaped detection  \\citep{Karlsson2008, Ekstrom2008d}. \\item Spectroscopic observations of halo stars \\citep[e.g.][]{Spite2005} indicate a primary production of nitrogen over a large metallicity range at low metallicity.  According to the chemical evolution models of \\citet{Chiappini2006, Chiappini2008}, primary nitrogen production by non-rotating or slow-rotating Pop III stars is not sufficient to explain the observations. For instance, primary nitrogen needs to be produced on larger ranges of masses and metallicities than expected in non rotating standard models. A similar plateau is observed in Damped Lyman Alpha (DLA) systems \\citep{Pettini2008}. \\item  Halo stars with log(O/H)+12 inferior to about 6.5 present higher C/O ratios than halo stars with log(O/H)+12 between 6.5 and 8.2 \\citep{Akerman2004, Spite2005}. Again a similar trend is observed in DLAs  \\citep{Pettini2008}. This is not predicted by slow rotating models. \\item The features listed above concern the bulk of the halo field stars. Now in addition to this ``normal'' population, there exists a group of stars showing a very different composition. These stars are collectively called Carbon-Enhanced Metal Poor stars (CEMP) and present high [C/Fe] ratios  \\citep[see the review by][]{Beers2005}.  The scatter in [C/Fe] is quite important indicating that these stars were not formed from a well mixed reservoir but acquired their peculiar high carbon abundance from locally enriched material. Some of these stars show also strong overabundances (with respect to iron) of nitrogen, oxygen and other heavy elements. At first order, these stars do appear to be formed from material having been processed mainly by H- and He-burning processes. \\item Very interestingly, halo stars in clusters also present some surprising features. While part of the cluster population follows the same trend as the ``normal'' population of the field, another part presents very different chemical patterns, with for instance, strong depletions of oxygen and strong enhancements in sodium \\citep[see the review by][]{Gratton2004}. These stars do appear to be formed from material having been processed only by H burning. \\end{itemize} Valid models should provide explanations for all these features. Ideally they should also provide some predictions for not yet observed characteristics. In the following we discuss the possible role of stellar axial rotation and we show how it can deeply affect the nucleosynthetic outputs of stars both in the massive ($>$ 8 M$_\\odot$) and intermediate mass range (2 $<$ M/M$_\\odot$ $<$ 8) and how it can explain some of the above observed features. ",
        "conclusions": ""
    },
    "0910/0910.0855_arXiv.txt": {
        "abstract": "In this paper we report on observations of the CoRoT LRa1 field with the Berlin Exoplanet Search Telescope (BEST). The current paper is part of the series of papers describing the results of our stellar variability survey. BEST is a small aperture telescope with a wide field-of-view (FOV). It is dedicated to search for stellar variability within the target fields of the CoRoT space mission to aid in minimizing false-alarm rates and identify potential targets for additional science. The LRa1 field is CoRoT's third observed field and the second long run field located in the galactic anticenter direction. We observed the LRa1 stellar field on $23$ nights between November and March $2005/2006$. From $6099$ stars marked as variable, $39$ were classified as periodic variable stars and $27$ of them are within the CoRoT FOV. We also confirmed the variability for $4$ stars listed in GCVS catalog. ",
        "introduction": "In previous articles belonging to a series dedicated to the variability survey in the CoRoT observational fields with BEST we presented the results of our survey on variable stars in the CoRoT IRa1~\\cite{Kabath07} and LRc1~\\cite{Karoff07} observational fields. The fields were observed with BEST~\\cite{Rauer04}, a small aperture telescope with a wide field-of-view (FOV) developed and operated by the Institut f\\\"{u}r Planetenforschung of Deutsches Zentrum f\\\"ur Luft- und Raumfahrt (DLR). As the BEST magnitude range covers approximately that of the CoRoT it is thus well suited for ground based support of the CoRoT space mission~\\cite{Baglin98}.  CoRoT was launched in December $2006$. The scientific goals of the mission are observations of selected stellar fields in order to find transiting extrasolar planets and to perform astroseismology of selected stars. The duration of the mission is planned to be $2.5$ years with $4$ long ($150$ days), $4$ short (up to $30$ days) and one initial ($30$ days) run. The CCD camera of CoRoT is divided into four segments from which two are dedicated to a transiting extrasolar planets survey and the other two to the asteroseismic survey. The total FOV covers $2.7^\\circ \\times 3.05^\\circ$. The point spread function (PSF) of CoRoT's extrasolar planet survey is about $80$ pixels for a $13$ mag K2 star and the PSF of the astroseismology field is about $410$ pixels for a solar type star at magnitude $5.7$ see~\\cite{Boisnard2006}. The prisms mounted in front of CoRoT's CCD provide colour information on observed objects which can help to determine the type of variability of the central star.  The BEST survey telescope system is used to discover and characterize variable stars in the CoRoT fields prior to CoRoT observations. The observational data can also be used as a complimentary information to the incoming CoRoT data. In addition BEST observations should also point out potentially interesting objects for the CoRoT additional science programs, such as binaries, $\\delta$ Scuti stars etc.. The advantage of our observations is the long time line which can provide extended lightcurves for several thousands of stars in the comparable magnitude range as CoRoT data and thus better understanding of e.g. potential variation of the lightcurve.  The resulting information about the variability in the stellar fields observed by BEST will be provided to the scientific community via the BEST archive linked to CoRoT's EXODAT~\\cite{Deleuil2006} database. ",
        "conclusions": "We performed photometry on CoRoT LRa1 field with BEST telescope to detect stellar variability. We identified $6099$ stars which were marked as variable and $44$ of them are showing a regular period. In total we detected $39$ new variable stars. The period ranges are usually between $0.1$ $<$ P $<$ $3$ days, however two new longperiodic stars with periods about $7$ and $21$ days were detected. The relatively small period range is given due to limited data set and due to duty cylce coverage which was disturbed by the bad weather period in December. We also confirm the period for the stars GU Mon, DD Mon and V 404 Mon which are already listed in the GCVS catalogue. The period for the GCVS star V 501 Mon could not be confirmed because no maxima/minima are present in our data set for that star. The star CD Mon is a longperiodic Mira type star and is also present in our data set. The newly found variable stars are currently observed within CoRoT's additional science programs. The information provided with our survey brings an additional information when analysing the CoRoT data and may avoid a false alarm for CoRoT transiting planet candidates. We gladly provide additional information about our data archive upon a request."
    },
    "0910/0910.5641_arXiv.txt": {
        "abstract": "{Realistic stellar models are essential to the forward modelling approach in asteroseismology. For practicality however, certain model assumptions are also required. For example, in the case of subdwarf B stars, one usually starts with zero-age horizontal branch structures without following the progenitor evolution. } {We analyse the effects of common assumptions in subdwarf B models on the $g$-mode pulsational properties. We investigate if and how the pulsation periods are affected by the H-profile in the core-envelope transition zone. Furthermore, the effects of C-production and convective mixing during the core helium flash are evaluated. Finally, we reanalyse the effects of stellar opacities on the mode excitation in subdwarf B stars.} {We computed detailed stellar evolutionary models of subdwarf B stars, and their non-adiabatic pulsational properties. Atomic diffusion of H and He is included consistently during the evolution calculations. The number fractions of Fe and Ni are gradually increased by up to a factor of 10 around $\\log T=5.3$. This is necessary for mode excitation and to approximate the resulting effects of radiative levitation. We performed a pulsational stability analysis on a grid of subdwarf B models constructed with OPAL and OP opacities.} {We find that helium settling causes a shift in the theoretical blue edge of the $g$-mode instability domain to higher effective temperatures. This results in a closer match to the observed instability strip of long-period sdB pulsators, particularly for $ l\\leq3$ modes. We show further that the $g$-mode spectrum is extremely sensitive to the H-profile in the core-envelope transition zone. If atomic diffusion is efficient, details of the initial shape of the profile become less important in the course of evolution. Diffusion broadens the chemical gradients, and results in less effective mode trapping and different pulsation periods. Furthermore, we report on the possible consequences of the He-flash for the $g$-modes. The outer edge of a flash-induced convective region introduces an additional chemical transition in the stellar models, and the corresponding spike in the Br\\\"unt-V\\\"ais\\\"al\\\"a frequency produces a complicated mode trapping signature in the period spacings. } {} ",
        "introduction": "Hot B-type subdwarfs are identified as core He-burning stars surrounded by a very thin H-envelope \\citep{heber1986}. This places them at the blue extension of the horizontal branch in the Hertzsprung-Russell diagram. Hence, they are also referred to as extreme horizontal branch (EHB) stars. Subdwarf B (sdB) stars are ubiquitous in our Galaxy, where they dominate the population of faint blue objects at high galactic latitudes \\citep{green1986}. They are also believed to be responsible for the ultraviolet excess or `UV upturn' in giant elliptical galaxies \\citep{brown1997,yi1997}.  A thorough review of hot subdwarfs is given by \\citet{heber2009}.  The future evolution of sdB stars is straightforward and undisputed. After He in the core is exhausted, a short phase of He-shell burning follows, and the star may be identified as a hotter subdwarf of O-type (sdO). The H-envelope is too thin to sustain H-shell burning, so the star will end as a C-O white dwarf without evolving through the asymptotic giant branch phase. The formation of sdBs on the other hand, remains a much debated topic. Many formation channels have been proposed such as enhanced mass loss on the RGB \\citep{d'cruz1996}, mass loss through binary interaction \\citep{mengel1976}, and the mergers of two He white dwarfs \\citep{webbink1984}. The relative importance of different formation scenarios has been evaluated by binary population synthesis studies \\citep{han2002,han2003}. While these studies are valuable, it should not be forgotten that they are dependent on parametric descriptions of binary interaction. Detailed mass determinations of subdwarf B stars could provide important constraints on the binary interaction mechanism. Accurate astrophysical mass determinations are only possible in special circumstances such as in eclipsing binary systems. Asteroseismology provides an alternative method, since it allows a detailed study of the interior of non-radially pulsating stars. The consistency of both approaches has been achieved for \\object{PG\\,1336$-$018} (see \\citealt{vuckovic2007}, \\citealt{hu2007}, and \\citealt{charpinet2008}).  Subdwarf B stars exhibit a variety of pulsations. The first sdB pulsator, \\object{EC\\,14026$-$2647}, was observed by \\citet{kilkenny1997} to pulsate in multiple short-period modes. This prototype represents the variable class  \\object{V361\\,Hya} stars which now includes 42 rapid sdB pulsators, with periods in the range \\mbox{80-600 s}. The short-period modes have been interpreted as low radial order, low spherical degree $p$-modes \\citep{charpinet1997}.  The \\object{V361\\,Hya} stars are amongst the hottest sdB stars with effective temperatures between 28 000 K and 35 000 K and surface gravities $5.2<\\log g<6.1$. At the cooler end of the EHB, 31 sdB pulsators with periods between 30 min and two hours have been discovered by \\citet{green2003}. This variable class is named V1093 Her but is also often referred to as \\object{PG\\,1716$+$426} stars after the class prototype. The long periods suggest that these stars pulsate in high radial order $g$-modes. Especially interesting are the sdB pulsators that exhibit both $p$- and $g$-mode pulsations. Three of these so-called hybrid pulsators have been found at the intersection of the \\object{V361\\,Hya} and \\object{V1093\\,Her} stars  \\citep{oreiro2005,baran2005,schuh2006,lutz2008}.  The opacity mechanism operating in the Fe opacity bump around $\\log T=5.3$ has been successfully used to explain the excitation of $p$- as well as the $g$-modes \\citep{charpinet1996,fontaine2003}. This opacity bump is caused by Fe accumulation owing to the competing diffusion processes of gravitational settling and radiative levitation. While seismic mass determinations have been achieved for a dozen \\object{V361\\,Hya} stars using static envelope models (see \\citealt{fontaine2008} and references therein), this is not yet the case for \\object{V1093\\,Her} stars. The reason is twofold. First of all, the $g$-modes have lower amplitudes and longer pulsation periods and thus require higher precision observations and longer observing runs to detect the frequencies. Secondly, envelope models may suffice to describe the \\object{V361\\,Hya} stars' shallow $p$-modes that probe only the outer layers. However, they are not suitable for modelling $g$-modes that penetrate to deeper stellar regions. Stellar models with detailed information about the He-rich core are required to predict reliable $g$-mode periods.  Another problem for the V1093 Her stars is the discrepancy between the observed and theoretically predicted instability strip. The first theoretical models that show unstable $g$-mode pulsations \\citep{fontaine2003} have much lower effective temperatures than observed with a discrepancy of $\\sim$ 5000 K. \\citet{jeffery2006b} managed to place the blue edge of the instability strip within $\\sim$1000 K of the observed blue edge. They did this by using OP opacities (\\citealt{badnell2005}) rather than OPAL opacities (\\citealt{iglesias1996}), and by enhancing the Ni abundance in the envelope in addition to Fe. Thus, their work suggests that  sophisticated models including up-to-date opacities and diffusion of other Fe-group elements are required to address the instability issue properly.  An exploratory study of how the pulsational properties of sdB stars are related to their internal structure has been carried out by \\citet{charpinet2000}. They showed that the He-H transition zone between the He-rich core and H-rich envelope is responsible for trapping  and confining $g$-modes in sdB stars. This is similar to the well-known mode trapping phenomenon in compositionally stratified white dwarfs \\citep{winget1981}. Mode trapping results in deviations from the asymptotic constant period spacing, making the $g$-modes sensitive to the mass of the H-envelope. Furthermore, the C-O/He-transition region between the convective and the radiative part of the core will also influence the $g$-mode period distribution. This may be a new interesting way to follow the He-core evolution, which would also have implications in a broader context for horizontal branch stars.  Here, we shall explore the sensitivity of the $g$-modes to certain assumptions about the He-flash. This energetic event characterizes the start of He-ignition in degenerate cores of low-mass stars  ($<$2 M$_{\\odot}$). Most stellar evolution codes encounter numerical difficulties in calculating this phase, and even in successful cases uncertainties remain because of the approximative treatment of convection. Three-dimensional hydrodynamic calculations are necessary to investigate this enigmatic phase of stellar evolution (e.g., \\citealt{deupree1996,dearborn2006,mocak2008}), but these are expensive in computing time. When many models are needed, as in forward modelling, it is more practical to construct post-flash models without following the previous evolution in detail. One should then use the outcome of detailed studies to ensure that the stellar models are as realistic as possible. Most simulations indicate that the initial flash starts off-centre, followed by several mini-flashes moving towards the centre. While it is generally found that most of the inner region will be convectively mixed, there are uncertainties in the outer extent of the convective region and the amount of He-burning during the He-flash varying from 3\\% to 7\\% in mass fraction \\citep{piersanti2004,serenelli2005}. The edge of the He-flash-induced convective region will leave a chemical composition gradient, which we believe can produce additional mode trapping features.  We start with constructing zero age extreme horizontal branch (ZAEHB) models in Sect.~\\ref{zaehb}. We make an analytical fit to the H-profile in the core-envelope transition region and examine its influence on the $g$-mode periods (Sect.~\\ref{profile}). The effects of uncertainties in the He-flash treatment, i.e.~convective mixing and C-production, are explored in Sect.~\\ref{heflash}. In Sect.~\\ref{ehb}, we perform a stability analysis on a grid of evolutionary sdB structures during the core He-burning phase. We show the impact of the opacities and H/He-diffusion on the instability strip. The conclusions are summarized in Sect.~\\ref{conclusions}. ",
        "conclusions": "The precise details of chemical transitions in subdwarf B stars have a considerable influence on the period spacings of high order $g$-modes. This is a signature of mode trapping that occurs due to discontinuities in the BV-frequency. We have shown that the structure of mode trapping depends sensitively on the exact shape of the H-profile in the core-envelope transition zone. Thus, when constructing ZAEHB models without following the previous evolution, the chemical profiles must be modelled carefully. Since the action of mode trapping is so sensitive to the shape of the chemical transition, it is necessary to include atomic diffusion during the evolution, i.e., the processes of gravitational settling, temperature and concentration diffusion.  Furthermore, if the sdB star started He-fusion in a run-away flash, the inner part of the core is convectively mixed and C is produced up to $\\sim$7\\%.  The edge of the He-flash-induced convective core is accompanied by a chemical transition.  We showed that this can cause additional mode trapping features resulting in a complicated behaviour of the period spacings. Multiple composition gradients lead to simultaneous mode trapping in different regions. The $g$-mode spectra become very complicated, and it will be difficult to directly infer detailed information about the composition gradients from observed $g$-modes. It is therefore crucial to evaluate the importance of different physical processes, and to account for all of them realistically, as we do here.  Finally, and perhaps most importantly, we report on a possible step forward in solving the `blue edge problem' of the V1093 Her instability strip. Our evolutionary sdB models with (i) OP opacities, (ii) the inclusion of H/He  diffusion, and (iii) a parametric Fe/Ni enhancement, show unstable $l\\leq 3$ $g$-modes for effective temperatures up to $30,000$ K. This is only achieved when all three ingredients are included in the computations. It should be noted that it remains unclear which $l$-values the observed long-period pulsation modes have. Thus, although we have demonstrated the importance of H/He diffusion, solving the blue edge problem effectively will require mode identification, in addition to a realistic treatment of radiative levitation and other transport mechanisms."
    },
    "0910/0910.3701_arXiv.txt": {
        "abstract": "Considering the high-energy limit of the QCD gluon distribution inside a nucleus, we calculate the proton-nucleus total inelastic cross section using a simplified dipole model. We show that, if gluon saturation occurs in the nuclear surface region, the total cross section of proton-nucleus collisions increases more rapidly as a function of the incident energy compared to that of a Glauber-type estimate. We discuss the implications of this with respect to recent ultra-high-energy cosmic ray experiments. ",
        "introduction": "In the ultra-high-energy domain, the mechanism of inelastic hadronic collisions is dominated by the contribution from small-$x$ gluons. This is the basic reason why the total or inelastic proton-proton collision cross section increases as a function of the incident energy. A simple way to see this is to use the eikonal expression for the reaction cross section in the impact parameter representation as% \\begin{equation} \\sigma _{r}=\\int d^{2}\\mathbf{b}[1-\\exp (-2\\chi (\\mathbf{b},s))], \\label{inela} \\end{equation}% where the eikonal $\\chi (\\mathbf{b},s)$ counts essentially the total number of possible scattering centers of the constituents inside the target \\textquotedblleft seen\\textquotedblright\\ by the projectile passing through the target along a straight line with the impact parameter $\\mathbf{b,}$ and with center-of-mass energy $\\sqrt{s}$.  If the target is thick enough, the eikonal function is much larger than unity ( $\\chi (\\mathbf{b},s)\\gg 1$) in the central region, but it falls down to zero in the surface region. The integrand of Eq.(\\ref{inela}) keeps a value almost equal to unity while $\\chi (\\mathbf{b},s)\\gg 1,$ and falls down quickly to zero near the surface. From this we may determine an effective radius $b_{1/2}$ such that the integral (\\ref{inela}) is approximately given by \\begin{equation} \\sigma _{r}\\simeq \\pi b_{1/2}^{2}. \\end{equation}% Here $b_{1/2}$ might be estimated by fixing the value of $\\chi (\\mathbf{b}% _{1/2},s)$, for example, as \\begin{equation} \\chi (\\mathbf{b}_{1/2},s)=\\ln 2, \\end{equation}% so that the effective radius depends on the energy $b_{1/2}=b_{1/2}\\left( \\sqrt{s}\\right) $. Let us assume that the eikonal function factorizes in the form  \\begin{equation} \\chi (\\mathbf{b},s)=P(\\mathbf{b})N(\\sqrt{s}), \\end{equation}% where $P(\\mathbf{b})$ is the probability distribution of the scattering center (e.g., partons) in the transverse plane and $N(x)$ is the number of partons which can interact for a given $\\sqrt{s}$. Suppose that $P(\\mathbf{b}% )$ is a two-dimensional Gaussian distribution of width $R,$ \\begin{equation} P(\\mathbf{b})=\\frac{1}{\\left( \\pi R\\right) ^{2}}e^{-\\frac{b^{2}}{R^{2}}} \\end{equation}% and that, for large $\\sqrt{s}$, the number of partons increase with $\\sqrt{s} $ as \\begin{equation} N(x)\\propto \\sqrt{s}^{\\alpha }. \\end{equation}% Then, we have \\begin{equation} b_{1/2}^{2}\\left( \\sqrt{s}\\right) =\\alpha R^{2}\\ln \\sqrt{s}+Const. \\end{equation}% In such a situation, the reaction cross section increases as function of the incident energy for very large $\\sqrt{s}$ as \\begin{equation} \\sigma _{r}\\simeq \\alpha \\pi R^{2}\\ln \\sqrt{s}.  \\label{sig1} \\end{equation}  If the edge of the distribution has an exponential tail instead of a Gaussian distribution, a similar argument will show that the cross section would increase as \\begin{equation} \\sigma _{r}\\simeq Const\\times \\left( \\ln \\sqrt{s}\\right) ^{2}.  \\label{sig2} \\end{equation}  The important point of this simple argument is that the rate of increase is related to the diffuseness $R$ of the probability distribution of the scattering center. The more diffuse the surface thickness is, the more quickly the reaction cross section increases as function of the incident energy. The possibility of such a mechanism, not only in the proton-proton but also in nucleus-nucleus cross section, was suggested many years ago~\\cite% {Kodama,Kodama1}.  In this paper we explore the idea of~\\cite{Kodama,Kodama1} in the language of the QCD gluon saturation mechanism for the proton-nucleus reaction. We show that, if the gluon distribution becomes saturated at some energy scale inside the nuclear surface region, then the reaction cross section of proton-nucleus collisions starts to increase very quickly and eventually overcomes the values estimated by the usual Glauber type of calculation.  Applying a simple effective dipole model for the reaction mechanism, we find that such an energy scale is of the order of $10^{17}-10^{18}eV.$ Above this energy scale, the behavior of the proton-nucleus cross section begins to change. We suggest that such an energy dependence of the proton-nucleus cross section may be observed in terms of the quantity called $\\left\\langle X_{\\max }\\right\\rangle $ in the air showers of ultra-high-energy cosmic rays. Using a very simple toy model estimate of $\\left\\langle X_{\\max }\\right\\rangle $, we show that our calculated values of the proton-nucleus reaction cross section are consistent with the recently observed $% \\left\\langle X_{\\max }\\right\\rangle $ by the Pierre Auger Observatory experiments~\\cite{Auger_Xmax} for incident protons at ultra-high energies. ",
        "conclusions": "In this paper, we have explored the idea of gluon saturation inside a nucleus in the high-energy limit and its effect on the proton-nucleus cross section. We show that if gluon saturation occurs in the surface region of the target nucleus, the proton-nucleus cross section starts to increase very rapidly as a function of the incident energy. Such a mechanism should eventually happen for an ultra-high energy scale, but the question is in which energy scale such a scenario starts to dominate.  Such energy scale is determined by the form of the distribution of gluons near the surface area. If we assume that the small $x$ gluon distribution in the nucleus follows that of the nucleon wave function inside the nucleus, the energy scale where the gluon saturation scenario starts to dominate the independent nucleon picture is around $10^{17}-10^{18}$ eV. We note that different profile functions give more or less the same energy scale once the proton-proton cross section is well fitted. It is very suggestive that the gluon saturation scenario inside the surface area seems consistent with the proton primary interpretation of the observed $\\left\\langle X_{\\max }\\right\\rangle $ behavior in the Pierre Auger Laboratory experiment.  As we see, the difference between the two pictures, INGP and GSNS, becomes very large at high energies. The reason for this is that, while in the INGP the effect of virtual gluons which bound the nucleon near the surface area is completely neglected, these gluons become dominant at high energies in the GSNS scenario. In a simple-minded argument, one might think that such an effect of nuclear binding must be negligible at high energies, since the ratio of the binding energy of a nucleon to the incident energy tends to zero. However, the situation may not be so simple. In the GSNS scenario, the density of virtual gluons, probably forming a kind of fractal fingers when penetrating into the vacuum among nucleons similar to the electric discharge pattern, becomes higher and higher at high energies, and eventually percolate everywhere even in the nuclear surface region. According to the color glass condensate picture~\\cite{Larry,Larry2,Larry3,Larry4}, such a scenario should happen at some energy scale, even at the lowest density region of the nuclear surface.  Naturally, the energy scale depends on the precise form of the geometric gluon distribution inside the nucleus. If the distribution does not follow the probability distribution of nucleons but more sharp surface distribution, then the energy scale may shift to higher energy. In the example of Fig. 3 and Fig. 4 the surface thickness parameter of the Woods-Saxon distribution is taken a little bit smaller than the usual value fitted to the nuclear distribution in nucleus ($d\\approx 0.6~fm$) since this fit only applies for heavy nuclei ($A>40$). If we take $d=0.6~fm,$ the energy scale is lower by one order of magnitude. Therefore, the energy scale depends crucially on how the gluons starts to saturate in the nuclear surface region. Depending on this, the energy scale can be even lower than estimated here. To find a real energy scale where the gluon saturation occurs at the nuclear surface, further investigation on high-energy proton-nucleus or electron-nucleus collisions will be necessary.  There exist ambiguities in the proton-proton cross section used in this paper (those of the model and extrapolation of the PDF, as well as those of experimental data) and these affect the precise value of the energy scale. However, the general conclusion of the present work does not change, as far as the energy dependence of the proton-proton cross section is fixed."
    },
    "0910/0910.4910_arXiv.txt": {
        "abstract": "{The identification and study of the first galaxies remains one of the most exciting topics in observational cosmology. The determination of the best possible observing strategies is a very important choice in order to build up a representative sample of spectroscopically confirmed sources at high-z ($z\\gtapprox7$), beyond the limits of present-day observations. } {This paper is intended to precisely adress the relative efficiency of lensing and blank fields in the identification and study of galaxies at $6\\ltapprox z \\ltapprox 12$. } {The detection efficiency and field-to-field variance are estimated from direct simulations of both blank and lensing fields observations. Present known luminosity functions in the UV are used to determine the expected distribution and properties of distant samples at $z\\gtapprox6$ for a variety of survey configurations. Different models for well known lensing clusters are used to simulate in details the magnification and dilution effects on the backgound distant population of galaxies. } {The presence of a strong-lensing cluster along the line of sight has a dramatic effect on the number of observed sources, with a positive magnification bias in typical ground-based ``shallow'' surveys ($AB\\ltapprox25.5$). The positive magnification bias increases with the redshift of sources and decreases with both depth of the survey and the size of the surveyed area. The maximum efficiency is reached for lensing clusters at $z\\sim0.1-0.3$. Observing blank fields in shallow surveys is particularly inefficient as compared to lensing fields if the UV LF for LBGs is strongly evolving at $z\\gtapprox7$. Also in this case, the number of $z\\ge8$ sources expected at the typical depth of JWST ($AB\\sim28-29$) is much higher in lensing than in blank fields (e.g. a factor of $\\sim10$ for $AB\\ltapprox28$). {All these results have been obtained assuming that number counts derived in clusters are not dominated by sources below the limiting surface brightness of observations, which in turn depends on the reliability of the usual scalings applied to the size of high-z sources. } } {Blank field surveys with a large field of view are needed to prove the bright end of the LF at $z\\gtapprox6-7$, whereas lensing clusters are particularly useful for exploring the mid to faint end of the LF. } ",
        "introduction": " ",
        "conclusions": "Summary and conclusions}  We have evaluated the relative efficiency of lensing clusters with respect to blank fields in the identification and study of $z\\ge6$ galaxies. The main conclusions of this study are given below.  For magnitude-limited samples of LBGs at $z\\ge6$, the magnification bias increases with the redshift of sources and decreases with both the depth of the survey and the size of the surveyed area. Given the typical near-IR FOV in lensing fields, the maximum efficiency is reached for clusters at $z\\sim0.1-0.3$, with maximum cluster-to-cluster differences ranging between 30 and 50\\% in number counts, depending on the redshift of sources and the LF.  The relative efficiency of lensing with respect to blank fields strongly depends on the shape of the LF, for a given photometric depth and FOV. The comparison between lensing and blank field number counts is likely to yield strong constraints on the LF.  The presence of a strong-lensing cluster along the line of sight has a dramatic effect on the observed number of sources, with a positive magnification effect in typical ground-based ``shallow'' surveys ($AB\\le25.5$). The postive magnification bias increases with the redshift of sources, and also from optimistic to pessimistic values of the LF. In case of a strongly evolving LF at $z\\ge7$, as proposed by \\cite{Bouwens08}, blank fields are particularly inefficient as compared to lensing fields. For instance, the size of the surveyed area in ground-based observations would need to increase by a factor of $\\sim10$ in blank fields with respect to a typical $\\sim30-40$ $arcmin^{2}$ survey in a lensing field, in order to reach the same number of detections at $z\\sim6-8$, and this merit factor increases with redshift. All these results have been obtained assuming that number counts derived in clusters are not dominated by sources below the limiting surface brightness of observations, which in turn depends on the reliability of the usual scalings applied to the size of high-z sources.  Ground-based ``shallow'' surveys are dominated by field-to-field variance reaching $\\sim$ 30 to 50\\% in number counts between z$\\sim$6 and 8 in a unique $\\sim30-40$ $arcmin^{2}$ lensing field survey (or in a 400 $arcmin^{2}$ blank field), assuming a strongly evolving LF.  The number of z$>$8 sources expected at the typical depth of JWST ($AB\\sim28-29$) is much higher in lensing than in blank fields if the UV LF is rapidly evolving with redshift (i.e. a factor of $\\sim10$ at $z\\sim10$ with $AB\\ltapprox28$).  Blank field surveys with a large FOV are needed to probe the bright edge of the LF at $z\\ge6-7$, whereas lensing clusters are particularly useful to explore the mid to faint end of the LF."
    },
    "0910/0910.4584_arXiv.txt": {
        "abstract": "{We determine the solar neutrino fluxes from a global analysis of the solar and terrestrial neutrino data in the framework of three-neutrino mixing. Using a Bayesian approach we reconstruct the posterior probability distribution function for the eight normalization parameters of the solar neutrino fluxes plus the relevant masses and mixing, with and without imposing the luminosity constraint. This is done by means of a Markov Chain Monte Carlo employing the Metropolis-Hastings algorithm. We also describe how these results can be applied to test the predictions of the Standard Solar Models. Our results show that, at present, both models with low and high metallicity can describe the data with good statistical agreement.}  \\preprint{YITP-SB-09-34\\\\IFT-UAM/CSIC-09-50\\\\EURONU-WP6-09-11} ",
        "introduction": "The idea that the Sun generates power through nuclear fusion in its core was first suggested in 1919 by Sir Arthur Eddington who pointed out that the nuclear energy stored in the Sun could explain the apparent age of the Solar System.  In 1939, Hans Bethe described in an epochal paper~\\cite{Bethe:1939bt} two nuclear fusion mechanisms by which main sequence stars like the Sun could produce the energy necessary to power their observed luminosities.  The two mechanisms have become known as the pp-chain and the CNO-cycle~\\cite{Bahcall:1989ks}.  For both chains the basic energy source is the burning of four protons to form an alpha particle, two positrons, and two neutrinos.  In the pp-chain, fusion reactions among elements lighter than $A = 8$ produce a characteristic set of neutrino fluxes, whose spectral energy shapes are known but whose fluxes must be calculated with a detailed solar model.  In the CNO-cycle the abundance of \\Nuc[12]{C} plus \\Nuc[13]{N} acts as a catalyst, while the \\Nuc[13]{N} and \\Nuc[15]{O} beta decays provide the primary source of neutrinos.  In order to precisely determine the rates of the different reactions in the two chains, which are responsible for the final neutrino fluxes and their energy spectrum, a detailed knowledge of the Sun and its evolution is needed. Standard Solar Models (SSM's)~\\cite{Bahcall:1987jc, TurckChieze:1988tj, Bahcall:1992hn, Bahcall:1995bt, Bahcall:2000nu, Bahcall:2004pz, PenaGaray:2008qe} describe the properties of the Sun and its evolution after entering the main sequence.  The models are based on a set of observational parameters (the present surface abundances of heavy elements and surface luminosity of the Sun, as well as its age, radius and mass) and on several basic assumptions: spherical symmetry, hydrostatic and thermal equilibrium, and equation of state.  Over the past five decades the solar models were steadily refined as the result of increased observational and experimental information about the input parameters (such as nuclear reaction rates and the surface abundances of different elements), more accurate calculations of constituent quantities (such as radiative opacity and equation of state), the inclusion of new physical effects (such as element diffusion) and the development of faster computers and more precise stellar evolution codes.  Despite the progress of the theory, only neutrinos, with their extremely small interaction cross sections, can enable us to see into the interior of a star and thus verify directly our understanding of the Sun~\\cite{Bahcall:1964gx}.  Indeed from the earliest days of solar neutrino research this test has been a primary goal of solar neutrino experiments, but for many years the task was made difficult by the increasing discrepancy between the predictions of the SSM's and the solar neutrino observations.  This so-called ``solar neutrino problem''~\\cite{Bahcall:1968hc, Bahcall:1976zz} was finally solved by the modification of the Standard Model of Particle Physics with the inclusion of neutrino masses and mixing. In this new framework leptonic flavors are no longer symmetries of Nature, and neutrinos can change their flavor from the production point in the Sun to their detection on the Earth. This flavor transition probability is energy dependent~\\cite{Pontecorvo:1967fh, Gribov:1968kq, Wolfenstein:1977ue, Mikheev:1986gs}, which explains the apparent disagreement among experiments with different energy windows. This mechanism is known as the LMA-MSW solution to the solar neutrino problem, and affects both the overall number of events in solar neutrino experiments and the relative contribution expected from the different components of the solar neutrino spectrum. Due to these complications, at first it was necessary to assume the SSM predictions for all the solar neutrino fluxes and their uncertainties in order to extract reasonably constrained values for neutrino masses and mixing. The upcoming of the real-time experiments Super-Kamiokande and SNO and the independent determination of the flavor oscillation probability using reactor antineutrinos at KamLAND opened up the possibility of extracting the solar neutrino fluxes and their uncertainties directly from the data~\\cite{Garzelli:2001ju, Bahcall:2002zh, Bahcall:2002jt, Bahcall:2004ut, Bahcall:2003ce, Bandyopadhyay:2006jn, PenaGaray:2008qe}. Nevertheless, in these works some set of simplifying assumptions had to be imposed in order to reduce the number of free parameters to be determined.  In parallel to the increased precision of the SSM-independent determination of the neutrino flavor parameters, a new puzzle has emerged in the consistency of SSM's~\\cite{Bahcall:2004yr}.  Till recently SSM's have had notable successes in predicting other observations. In particular, quantities measured by helioseismology such as the radial distributions of sound speed and density~\\cite{Bahcall:1992hn, Bahcall:1995bt, Bahcall:2000nu, Bahcall:2004pz} showed good agreement with the predictions of the SSM calculations and provided accurate information on the solar interior.  A key element to this agreement is the input value of the abundances of heavy elements on the surface of the Sun~\\cite{Grevesse:1998bj}. However, recent determinations of these abundances point towards substantially lower values than previously expected~\\cite{Asplund:2004eu, Asplund:2009fu}. A SSM which incorporates such lower metallicities fails at explaining the helioseismological observations~\\cite{Bahcall:2004yr}, and changes in the Sun modeling (in particular of the less known convective zone) are not able to account for this discrepancy~\\cite{Chaplin:2007uh, Basu:2006vh}.  So far there has not been a successful solution of this puzzle. Thus the situation is that, at present, there is no fully consistent SSM. This led to the construction of two different sets of SSM's, one (labeled ``GS'') based on the older solar abundances~\\cite{Grevesse:1998bj} implying high metallicity, and one (labeled ``AGS'') assuming lower metallicity as inferred from more recent determinations of the solar abundances~\\cite{Asplund:2004eu, Asplund:2009fu}. In Ref.~\\cite{PenaGaray:2008qe} the solar fluxes corresponding to such two models were detailed, based on updated versions of the solar model calculations presented in Ref.~\\cite{Bahcall:2004pz}. These fluxes were denoted as ``BPS08(GS)'' and ``BPS08(AGS)'', respectively. In a very recent work~\\cite{Serenelli:2009yc} an update of the BPS08(AGS) solar model has been constructed using the latest determination of the compositions~\\cite{Asplund:2009fu} as well as some improvement in the equation of state. For what concerns the overall normalization of solar neutrino fluxes, the predictions of this new model are very close to those of BPS08(AGS).  In this work we perform a solar model independent analysis of the solar and terrestrial neutrino data in the framework of three-neutrino masses and mixing. The aim of this analysis is to simultaneously determine the flavor parameters and all the solar neutrino fluxes with a minimum set of theoretical priors. In Sec.~\\ref{sec:ana} we present the method employed, the data included in the analysis and the physical assumptions used in this study.  The results of the analysis are given in Sec.~\\ref{sec:res}, where we show the reconstructed posterior probability distribution function for the eight normalization parameters of the solar neutrino fluxes.  We discuss in detail the effect of the luminosity constraint~\\cite{Bahcall:2001pf} as well as the role of the Borexino experiment and its potential for improvement.  In addition, we use the results of this analysis to statistically test to what degree the present solar neutrino data can discriminate between the two SSM's.  Finally in Sec.~\\ref{sec:sum} we summarize our conclusions. ",
        "conclusions": "\\label{sec:sum}  We have performed a solar model independent analysis of the solar and terrestrial neutrino data in the framework of three-neutrino oscillations, following a Bayesian approach in terms of a Markov Chain Monte Carlo using the Metropolis-Hastings algorithm.  This approach has allowed us to reconstruct the probability distribution function in the entire eleven-dimensional parameter space, consistently incorporating the required set of theoretical priors. The best fit values and allowed ranges for the eight solar neutrino fluxes are summarized in Eq.~\\eqref{eq:bestlc} and Fig.~\\ref{fig:fitlc} for the analysis with the luminosity constraint, and in Eq.~\\eqref{eq:bestnolc} and Fig.~\\ref{fig:fitnolc} for the more general case of unconstrained solar luminosity. We found that at present the neutrino-inferred luminosity perfectly agrees with the measured one and it is known with a $1\\sigma$ uncertainty of 14\\%. We have also tested the fractional energy production in the pp-chain and the CNO-cycle, finding that the total amount of the solar luminosity produced in the CNO-cycle is bounded to be $L_\\text{CNO} / L_\\odot < 2.8\\%$ at 99\\% CL  irrespective of whether the luminosity constraint is imposed or not.  We have then presented a statistical test which can be performed with these results in order to shed some light on the so-called solar composition problem, which at present arises in the construction of the Standard Solar Model.  We found that the low value of the \\Nuc[8]{B} flux measured at SK and SNO points towards low metallicity models, whereas the measurement of \\Nuc[7]{Be} in Borexino and the corresponding value of the \\Nuc{pp} flux implied by the luminosity constraint show better agreement with high metallicity models. Altogether the fit shows a slight preference for models with higher metallicities, however the difference between the two models is not very significant at statistical level. While a realistic improvement of the Borexino data analysis in the near future can positively affect the direct determination of most solar neutrino fluxes, it is unlikely that enough precision will be achieved to go beyond the the present theoretical uncertainties of the SSM's. The largest difference between the models lies on the CNO fluxes that give predictions which differ by about 30\\%. Thus ideally in order to achieve a statistically meaningful discrimination between the models one would need a low energy solar neutrino experiment capable of measuring the neutrino energy spectrum for energies between $0.5~\\text{MeV} \\lesssim E_\\nu \\lesssim 1.5~\\text{MeV}$ and, more importantly, of rejecting the radioactive backgrounds to the required level, so to allow for a determination of the CNO fluxes at $\\mathcal{O}(\\text{30\\%})$ precision."
    },
    "0910/0910.4482.txt": {
        "abstract": "We present high precision radial velocities (RVs) of double-lined spectroscopic binary stars HD78418, HD123999, HD160922, HD200077 and HD210027. They were obtained based on the high resolution echelle spectra collected with the Keck I/Hires, Shane/CAT/Hamspec and TNG/Sarge telescopes/spectrographs over the years 2003-2008 as a part of TATOOINE search for circumbinary planets. The RVs were computed using our novel iodine cell technique for double-line binary stars which relies on tomographically disentangled spectra of the components of the binaries. The precision of the RVs is of the order of 1-10 m~s$^{-1}$ and to properly model such measurements one needs to account for the light time effect within the binary's orbit, the relativistic effects and the RV variations due to the tidal distortions of the components of the binaries. With such proper modeling, our RVs combined with the archival visibility measurements from the Palomar Testbed Interferometer allow us to derive very precise spectroscopic/astrometric orbital and physical parameters of the binaries. In particular, we derive the masses, the absolute K and H band magnitudes and the parallaxes. The masses together with the absolute magnitudes in the K and H bands enable us to estimate the ages of the binaries.  These RVs allow us to obtain some of the most accurate mass determinations of binary stars. The fractional accuracy in $m\\sin i$ only and hence based on the RVs alone ranges from 0.02\\% to 0.42\\%. When combined with the PTI astrometry, the fractional accuracy in the masses ranges in the three best cases from 0.06\\% to 0.5\\%.  Among them, the masses of HD210027 components rival in precision the mass determination of the components of the relativistic double pulsar system PSR~J0737-3039. In the near future, for double-lined eclipsing binary stars we expect to derive masses with a fractional accuracy of the order of up to $\\sim$0.001\\% with our technique. This level of precision is an order of magnitude higher than of the most accurate mass determination for a body outside the Solar System --- the double neutron star system PSR~B1913+16. ",
        "introduction": "The first observations of a spectroscopic binary star, $\\zeta$ UMa (Mizar), were announced by Edward C. Pickering (1846-1919) on 13 November 1889 during a meeting of the National Academy of Sciences in Philadelphia \\citep{Pickering:90::}. A similar announcement about $\\beta$ Per (Algol) was made by Herman C. Vogel (1841-1907) on 28 November 1889 during a session of Konglich-Preussiche Akademie der Wissenschaften \\citep{Vogel:90a::,Vogel:90b::}. Even though a transatlantic telegraph cable had been available since the 1860s, it was quite unusual timing for a pre-astro-ph era. Pickering noted that the K line of Mizar occasionally appeared double and this way discovered the first double-lined spectroscopic binary star. Vogel measured radial velocities (RVs) of Algol and used them to prove that the known variations in the brightness of Algol are indeed caused by a ``dark satellite revolving about it\" \\citep{Vogel:90b::}.  Vogel and collaborators built a series of prism spectrographs for the 30-cm refractor of the Potsdam Astrophysical Observatory \\citep{Vogel:00::}. In 1888, they initiated a photographic radial-velocity program. Vogel was not the first to take a photograph of a stellar spectrum but improved the technique and obtained an RV precision of 2-4 km~s$^{-1}$ \\citep{Vogel:91::}. Vogel (1906) and Pickering (1908) were both awarded the Bruce medal. Yet another Bruce medalist (1915), William W. Campbell (1862-1938) and collaborators carried out a large spectroscopic program at the Lick Observatory and discovered many spectroscopic binaries \\citep{Camp:11::}. They prepared the first catalog of spectroscopic binaries \\citep{Camp:05::}. Vogel was the director of the Potsdam Observatory for 25 years, Pickering of the Harvard College Observatory for 42 years and Campbell of the Lick Observatory for 30 years. The administrative duties must have been not too distracting those days.  Both Vogel and Campbell recognized the importance of flexure and temperature control of a spectrograph \\citep{Vogel:90c::,Camp:98::,Camp:00::} and Campbell also noted the issue of a slit illumination and its impact on RV precision \\citep{Camp:16::}. Today a high RV precision is achieved by either using a very stable fiber fed spectrograph that is contained in a controlled environment (a vacuum tank) or using an absorption cell to superimpose a reference spectrum onto a stellar spectrum and this way measure and account for the systematic RV errors. Current state of the art precision is at the level of $\\sim$1 m~s$^{-1}$. It is however important to note that such a precision refers to single stars or at best single-lined spectroscopic binaries where the influence of the secondary spectrum can be neglected. In such a case, given a stable spectrograph, an RV measurement is essentially a measurement of a shift of an otherwise constant shape (spectrum).  Radial velocities (RVs) of double-lined spectroscopic binary stars (SB2) can be used effectively to derive basic parameters of stars if the stars happen to be eclipsing or their astrometric relative orbit can be determined. It is quite surprising that the RV precision of double-lined binary stars on the average has not improved much over the last 100 years (see Fig.~\\ref{fig1}). With the exception of our previous work \\citep{Konacki:05::,Konacki:09a::}, the RV precision for such targets typically varies from $\\sim$0.1  km~s$^{-1}$ to $\\sim$1 km~s$^{-1}$ and clearly is much worse than what has been achieved for stars with planets or single-lined binary stars. The main problem with double-lined binary stars is that one has to deal with two sets of superimposed spectral lines whose corresponding radial velocities change considerably with typical amplitudes of $\\sim$ 50-100 km~s$^{-1}$. In consequence a spectrum is highly variable and obviously one cannot measure RVs by noting a simple shift.  We have developed a novel iodine cell based approach that employs a tomographic disentangling of the component spectra of SB2s and allows one to measure RVs of the components of SB2s with a precision of the order 1-10 m$\\,$s$^{-1}$ \\citep{Konacki:09a::,Konacki:09b::}. Such quality RVs not only enable us to search for circumbinary extrasolar planets \\citep{Konacki:09b::} but also to determine basic parameters of stars with an unprecedented precision. In particular the masses of stars for non eclipsing SB2s can be easily determined with a fractional accuracy of the order of at least $\\sim$$0.1\\%$ and often even $\\sim$$0.01\\%$. Moreover, we expect that the accuracy in masses will reach $\\sim$$0.001\\%$ level when our method is applied to eclipsing binary stars. Such a level of precision is an order of magnitude higher than of the most accurate mass determination for a body outside the Solar System --- the double neutron star system PSR~B1913+16 \\citep{Nice:08::}.  Below we present our precision RV data sets for five targets HD78418, HD123999, HD160922, HD200077 and HD210027 from our on-going TATOOINE (The Attempt To Observe Outer-planets In Non-single-stellar Environments) RV program to search for circumbinary planets. All of them have been extensively observed with the Palomar Testbed Interferometer \\citep[PTI;]{Col:99a::}. The archival PTI visibility measurements can be used to derive relative astrometric orbits of the binaries. These combined with our spectroscopic orbits allow for a complete orbital and physical description of the systems (with the exception of the radii of the components). In \\S{2} we describe the RV measurements and their modeling and in \\S{3} the visibility measurements and their modeling. In \\S{4} we present the spectroscopic and astrometric orbital solutions and the resulting orbital and physical parameters of the binaries. A discussion is provided in \\S{5}. ",
        "conclusions": "There are a few ways to determine accurate masses of normal stars. Perhaps the most classic one is through absolute astrometric orbits of both components of a binary. However this constitues a challenging measurement and in practice masses are typically derived using diverse data sets e.g. (1) by combining RVs, relative astrometry and parallax for single-lined spectroscopic binaries, (2) as in this paper by combining relative astrometry and RVs for double-lined spectroscopic binaries and (3) by combining light curves and RVs for eclipsing double-lined binaries. The last method is the most useful one as it not only provides the most accurate masses due to a convenient edge-on geometry (note that the masses are derived from their respective M $\\sin^3i$) but also enables one to determine the radii of stars.  In a recent review, \\cite{Torres:09::} collect 118 detached binary stars (including 94 eclipsing) with the most accurate mass determinations in the literature. These are denoted with open circles in Fig.~18. Even though our targets are not eclipsing and the orbital inclinations and their errors have a significant impact on the precision in masses, for two stars HD123999 and HD210027 we have obtained more accurate mass determinations than for any of the stars from \\cite{Torres:09::}. The accuracies of the masses of these two binaries are in the precision range covered by only close double neutron star systems characterized with radio pulsar timing. In particular, the masses for HD210027 rival in precision the mass determination of the components of the relativistic double pulsar system PSR~J0737-3039 \\citep{Nice:08::}. If our targets were all eclipsing and the accuracy in masses was limited by our RVs alone, the accuracy would be in the range 0.02\\% to 0.42\\% (the fractional accuracy of M $\\sin^3i$). The lower limit of this range is equal to the mass accuracy of PSR~B1913+16 which has the most accurate mass determination for a body outside the Solar System \\citep{Nice:08::}. In fact if we adopt our RV precision and use the achievable from the ground precision in the orbital inclination angle for eclipsing binaries, we can expect to obtain masses with a fractional precision of $\\sim$0.001\\% (see Fig.~18). Clearly, our RV technique for double-lined and eclipsing binary stars opens an exciting opportunity for derving masses of stars (and other parameters) with an unprecedented precision. These combined with parallax measurements from e.g. the planned GAIA astrometric mission and hopefully accurate abundance determinations should produce an outstanding set of parameters to tests models of the stellar structure and evolution."
    },
    "0910/0910.2641_arXiv.txt": {
        "abstract": "We present an overview on how variability can be used to constrain the location of the ionized outflow in nearby Active Galactic Nuclei using high-resolution X-ray spectroscopy. Without these constraints on the location of the outflow, the kinetic luminosity and mass loss rate can not be determined. We focus on the Seyfert 1 galaxy NGC 5548, which is arguably the best studied AGN on a timescale of 10 years. Our results show that frequent observations combined with long term monitoring, such as with the \\textit{Rossi X-ray Timing Explorer (RXTE)} satellite, are crucial to investigate the effects of these outflows on their surroundings. ",
        "introduction": "Active Galactic Nuclei (AGN) outflows are thought to play an important role in feedback processes. The growth of the supermassive black hole is connected to the growth of the bulge and the interstellar and intergalactic medium are thought to be enriched by these AGN outflows \\cite{DiMatteo05}. However the main problem is that we do not know the location or origin of these outflows. The kinetic luminosity and mass loss rate can not be accurately determined if the location is unknown. By studying these outflows in the X-ray regime with high-resolution grating spectrometers over multiple years, we can constrain the location of the outflow by using the intrinsic variability of the source \\cite{Netzer03, Nicastro07}. We will focus on one particular Seyfert 1 galaxy, NGC 5548, because it is the best studied AGN, with high-resolution X-ray data spanning almost 8 years in total \\cite{Detmers08}. ",
        "conclusions": "Long-term monitoring of AGN in combination with high-resolution X-ray spectroscopy is key to constrain the location of the warm absorber outflows. Once the location has been determined, the feedback strength of the outflow (mass outflow rate, kinetic luminosity) can not be accurately determined without knowing the geometry of the outflow. General estimates can be made, either by assuming that the mass outflow rate is less than the accretion rate or by assuming a specific geometry for the outflow, but the key to unravelling the importance of feedback in these local AGN is variability, combined with multi-wavelength observations to constrain the geometry of these outflows. Only then can a full picture of feedback in local AGN be obtained.  \\begin{theacknowledgments} RGD wishes to thank Elisa Costantini for useful discussions regarding the paper. SRON is supported financially by NWO, The Netherlands Organization for Scientific Research. \\end{theacknowledgments}"
    },
    "0910/0910.0716_arXiv.txt": {
        "abstract": "{The photometric, structural and kinematical  properties of the centers of elliptical galaxies, harbour important information of the formation history of the galaxies. In the case of non active elliptical galaxies these properties are linked in a way that surface brightness, break radius and velocity dispersion of the core lie on a fundamental plane similar to that found for their global properties.} {We construct the Core Fundamental Plane (CFP) for a sizeable sample of low redshift radio galaxies and compare it with that of non radio ellipticals.} {To pursue this aim we combine  data obtained from high resolution HST images with medium resolution optical spectroscopy to derive the photometric and kinematic properties of $\\sim$40 low redshift radio galaxies. } {We find that the CFPs of radio galaxies is indistinguishable from that defined by non radio elliptical galaxies of similar luminosity. The characteristics of the CFP of radio galaxies are also consistent (same slope) with those of the Fundamental Plane (FP) derived from the global  properties of radio  (and non radio) elliptical galaxies. The similarity of CFP and FP for radio and non radio ellipticals suggests that the active phase of these galaxies has minimal effects for the structure of the galaxies.} {} ",
        "introduction": "The properties of the centers of massive elliptical galaxies are of strategic importance for the understanding of the complex processes of galaxy formation. The centers represent the bottom of the potential well of the galaxy, host massive black holes and provide a record of the past history of the galaxies. Until a decade ago the properties of the center of galaxies were very little studied because of insufficient spatial resolution of available instrumentation. Only with the use of Hubble Space Telescope, and using future large ground based telescope such a study become feasible.   In the pioneering study on the nuclear properties of nearby early type galaxies (carried out with HST images) \\cite{faber} were able to probe the inner regions of a number of nearby ellipticals. They point out that the inner luminosity profiles can be parameterized by a core (or break) radius $R_b$ and a characteristic surface brightness within the  break radius $<\\mu_b>$. They also showed that these parameters ($R_b$, $<\\mu_b>$) can be combined to form a Core Fundamental Plane (CFP) which is analogous to the one found for the global properties of early type galaxies (\\cite{faber}). Assuming that the cores are in dynamical equilibrium and supported by random motions, and that the velocity anisotropy does not vary too much from galaxy to galaxy, and if the core $M/L$ is a well-behaved function of any two variables $<\\mu_b>$, $R_b$, or $\\sigma_0$,  then one expects that  the cores of galaxies follow a law similar to that of the Fundamental Plane (FP).  The main reason to study global  and core properties is that it allows one to probe the mechanism of galaxy formation. On the other end in the last decade a new ingredient in this picture is represented by the discovery of the presence of a massive black bole (BH) in the centers of virtually all galaxies (e.g. \\cite{FM}, \\cite{lauer}). However, only a small fraction of these BHs exhibit associated nuclear activity (non thermal nuclear emission, X-ray and radio emission). The BHs may play an important role in the formation and evolution of massive galaxies and are also a key component for the development of the nuclear activity. Therefore the comparison of the properties of active and inactive galaxies through the FP becomes a tool to investigate the interplay between galaxy formation and nuclear activity.  In a previous work, using photometrical and dynamical data for 73 low red-shift (z$<$0.2) radio galaxies (RG), we (\\cite{bettoni}) were able to compare the FP of RG  with the one defined by inactive ellipticals (\\cite{jfk96}, JFK96). We showed that the same FP holds for both radio and non radio ellipticals with radio galaxies occupying the region of the most luminous and large galaxies.  Till very recently few data were available on the optical nuclear properties of radio galaxies. One of the best optical sets of data is part of the study of the B2 sample of low luminosity radio galaxies (\\cite{fanti}, \\cite{Cap00}). For $\\sim$60 radio galaxies WFPC2 HST imaging is available in the $F555W$ (approximately $V$)  and $F814W$ (approximately $I$) filters. A similar set of data is also available for a sample of powerful 3C radio sources. These studies revealed the presence of new and interesting features, some of them almost exclusively associated to low luminosity FR I radio galaxies. In particular, the HST observations have shown the presence of dust in a large fraction of weak (FR I) radio galaxies which takes the form of extended nuclear disks (\\cite{jaffe}, \\cite{dekoff}, \\cite{dejuan}, \\cite{verdoes}, \\cite{hans}). Such structures have been naturally identified with the reservoir of material which will ultimately accrete into the central black hole.  The symmetry axis of the nuclear disk may be a useful indicator of the rotation axis of the central black hole (see e.g. \\cite{capetti1}), although the precise relationship between these two axes remains uncertain.  In this paper we aim to investigate the CFP for a sample of low redshift radio galaxies and to compare it with that for radio quiet galaxies. The plan of the paper is as follows. In Section 2 we discuss our observations and data reduction methods. In Section 3 we present the CFP for our sample of RG. The implications of those observations are then discussed in Section 4. Throughout this paper we use the Concordance Cosmological Model, with $H_0=70$ kms$^{-1}$Mpc$^{-1}$, and $\\Omega_{\\Lambda}=0.70$.   \\begin{table*} \\caption{The sample of low redshift B2 radio galaxies with velocity dispersion measurements} \\label{tab_b2} \\begin{center} \\begin{tabular}{lcccccccc } \\hline name & $m_V$ &    z  &    $\\sigma_c$ & $\\Delta$$\\sigma$ &  $<\\mu_b>$ &  $\\Delta$$<\\mu_b>$ & $R_b$  & $\\Delta$$r_b$ \\\\ (1) & (2) & (3) & (4) & (5) & (6) & (7) & (8) & (9) \\\\ \\hline 0648+27$^+$   &  13.82  &  0.0409 &  137.8  & 38.0 & 14.94 & 0.05 & 0.26 & 0.02  \\\\ 0755+37    &  13.80  &  0.0413 &  291.0  & 17.3 & 17.05 & 0.15 & 1.09 & 0.09  \\\\ 0908+37    &  15.67  &  0.1040 &  387.1  & 18.0 & 17.36 & 0.10 & 0.50 & 0.05  \\\\ 0915+32B &  15.20  & 0.0620 &  252.1  & 24.3 & 17.00 & 0.05 & 0.52 & 0.03 \\\\ 0924+30    &  13.32  & 0.0266 &  243.7  & 14.5 & 17.33 & 0.15 & 0.96 & 0.14 \\\\ 1003+26    &  15.47  & 0.1165 &  438.4  & 21.6 & 17.94 & 0.15 & 0.41 & 0.05 \\\\ 1113+24    &  15.10  & 0.1021 &  377.8  & 10.2 & 17.44 & 0.05 & 0.47 & 0.03 \\\\ 1204+34    &  15.35  & 0.0788 &  201.0  & 29.6 & 17.47 & 0.25 & 0.46 & 0.09 \\\\ 1322+36B &  12.84  & 0.0175 &  259.1  & 12.1 & 15.45 & 0.05 & 0.62 & 0.05 \\\\ 1339+26B  &  15.10 & 0.0757 &  379.1  & 18.2 & 16.41 & 0.15 & 0.32 & 0.04 \\\\ 1347+28    &  15.66  & 0.0724 &  228.0  & 19.4 & 17.18 & 0.05 & 0.34 & 0.03 \\\\ 1357+28    &  14.81  & 0.0629 &  308.3  & 11.2 & 16.80 & 0.05 & 0.40 & 0.03  \\\\ 1422+26B &  14.32  & 0.0370 &  264.9  & 11.6 & 15.44 & 0.05 & 0.13 & 0.01 \\\\ 1430+25    &  16.65  & 0.0813 &  203.6  & 14.3 & 16.72 & 0.25 & 0.13 & 0.03 \\\\ 1447+27    &  13.88  & 0.0306 &  318.4  & 11.6 & 15.77 & 0.15 & 0.47 & 0.04 \\\\ 1450+28  &  16.31  &  0.1203  & 405.6  & 23.6 & 16.18 & 0.05 & 0.16 & 0.15 \\\\ 1450+28*$^+$  &  16.31   & 0.1265 &  333.3  & 22.1 & 16.56 & 0.05 & 0.20 & 0.01 \\\\ 1502+26   &  15.42   & 0.0540 &  402.1  & 19.9 & 16.60 & 0.15 & 0.66 & 0.07 \\\\ 1512+30   &  15.37   & 0.0931 &  330.2  & 25.6 & 16.55 & 0.05 & 0.34 & 0.03 \\\\ 1521+28   &  15.09   & 0.0825 &  323.0  & 11.8 & 18.02 & 0.15 & 0.85 & 0.06 \\\\ 1527+30   &  15.64   & 0.1143 &  448.0  & 23.0 & 16.76 & 0.10 & 0.34 & 0.03 \\\\ 1553+24   &  14.41   & 0.0426 &  269.9  & 10.9 & 16.41 & 0.05 & 0.45 & 0.03 \\\\ 1557+26   &  14.97   & 0.0442 &  310.7  & 16.1 & 15.45 & 0.05 & 0.21 & 0.01 \\\\ 1613+27   &  14.90   & 0.0647 &  246.6  & 13.5 & 16.74 & 0.15 & 0.30 & 0.03 \\\\ 1658+30   &  14.47   & 0.0351 &  315.8  & 10.2 & 15.02 & 0.15 & 0.19 & 0.02 \\\\ 1726+31   &  17.04   & 0.1670 &  294.8  & 26.8 & 15.96 & 0.15 & 0.22 & 0.03 \\\\ 1827+32   &  15.10   & 0.0659  & 308.5  & 11.4 & 16.24 & 0.05 & 0.21 & 0.01 \\\\ 2236+35   &  12.57   & 0.02759 & 297.8   & 12.5 & 17.09 & 0.05 & 1.28 & 0.11  \\\\ \\hline 0034+25 &  13.35 &  0.031849   &   295.0   & 10.0 & 16.27 & 0.05 & 0.74 & 0.04 \\\\ 0055+26 &  13.80 &  0.047400   &   231.0   & 13.0 & 16.34 & 0.05 & 0.40 & 0.02 \\\\ 0120+33  & 11.43  &  0.016458   &   315.0   & 10.0 & 15.13 & 0.05 & 0.75 & 0.07 \\\\ 1217+29 &  10.41 &  0.002165   &   238.0   & 10.0 & 13.87 & 0.05 & 0.80 & 0.06 \\\\ 1256+28 &  13.88 &  0.022879   &   203.0   & 10.0 & 16.48 & 0.15 & 0.70 & 0.07 \\\\ 1257+28 &  11.94 &  0.024097   &   278.0   & 10.0 & 17.97 & 0.10 & 2.55 & 0.01 \\\\ 1525+29 &  14.95  &  0.065155   &   250.0   & 30.0 & 16.56 & 0.15 & 0.53 & 0.05 \\\\ 1610+29 &  13.29 &  0.031852   &   331.0   & 26.0 & 16.49 & 0.05 & 0.95 & 0.06 \\\\ 3C88       &  14.24  & 0.030221    &   190.0   & 22.0 & 17.47 & 0.05 & 0.98 & 0.09 \\\\ 3C192     &  15.72  & 0.059709    &   199.0   & 23.0 & 14.34 & 0.15 & 0.09 & 0.01 \\\\ 3C388     &  14.77 & 0.091700    &   365.0   & 23.0 & 18.03 & 0.15 & 0.96 & 0.09 \\\\ \\hline \\end{tabular} \\end{center} \\tiny{Columns: (1) identification of the source, (2) apparent magnitude $m_V$ from \\cite{col75} (photographic magnitudes converted to visual magnitudes, \\cite{fan78}, \\cite{gon93}, and \\cite{gon00}, (3) redshift, (4, 5) measured velocity dispersion and error in km/sec, (8 ,9), $<\\mu_b>$ and error in mag/$arcsec^2$, (8, 9) break radius $r_b$ and error in arcsecs,  from \\cite{hans1};~ *companion of the radio galaxy, + not used in the CFP see Appendix A1}  \\end{table*} ",
        "conclusions": "We have presented the photometric, structural and kinematic properties  of the centers of a sample of 38 low redshift radio galaxies galaxies and have shown that the conventional parameters characterizing the centers (the break radius, its surface brightness and the central velocity dispersion) are well represented in a plane (the Core Fundamental Plane) which is indistinguishable from that of elliptical galaxies that do not exhibit radio emission.  A similar result was found from the comparison of the  global properties of a sample of 72  radio galaxies and local Ellipticals using the standard Fundamental Plane description (\\cite{bettoni}). The comparison of the properties of the FP, that refer to the whole galaxy, with those concering the centers (the CFP) shows that the slopes of the two  planes are very similar and suggest taht the same mechanism is responsible for the link among the involved quantities.  The remarkable similarity of the properties of radio and non radio elliptical galaxies in the description of both the Fundamental Plane and the Core Fundamental Plane indicates that the active phase of the galaxy connected with the strong emission at radio frequencies has likely an inconsequential effect for the structure of the whole galaxy.  Moreover these results suggest that the two type of galaxies (radio and non radio) have had a similar history of formation and evolution.    \\appendix"
    },
    "0910/0910.5305_arXiv.txt": {
        "abstract": "We conducted radio detection observations at 8.4~GHz for 22~radio-loud broad absorption line~(BAL) quasars, selected from the Sloan Digital Sky Survey~(SDSS) Third Data Release, by a very-long-baseline interferometry~(VLBI) technique.  The VLBI instrument we used was developed by the Optically ConnecTed Array for VLBI Exploration project~(OCTAVE), which is operated as a subarray of the Japanese VLBI Network~(JVN).  We aimed at selecting BAL quasars with nonthermal jets suitable for measuring their orientation angles and ages by subsequent detailed VLBI imaging studies to evaluate two controversial issues of whether BAL quasars are viewed nearly edge-on, and of whether BAL quasars are in a short-lived evolutionary phase of quasar population.  We detected 20 out of 22 sources using the OCTAVE baselines, implying brightness temperatures greater than 10$^{5}$~K, which presumably come from nonthermal jets.  Hence, BAL outflows and nonthermal jets can be generated simultaneously in these central engines.  We also found four inverted-spectrum sources, which are interpreted as Doppler-beamed, pole-on-viewed relativistic jet sources or young radio sources: single edge-on geometry cannot describe all BAL quasars.  We discuss the implications of the OCTAVE observations for investigations for the orientation and evolutionary stage of BAL quasars. ",
        "introduction": "Broad absorption line (BAL) quasars are a subclass of active galactic nuclei~(AGNs) with rest-frame ultra-violet spectra showing absorption troughs displaced blueward from the corresponding emission lines in the high-ionization transitions of C$_\\mathrm{IV}$, Si$_\\mathrm{IV}$, N$_\\mathrm{V}$, and O$_\\mathrm{IV}$, and occasionally in low-ionization transitions, e.g., Mg$_\\mathrm{II}$ and Al$_\\mathrm{III}$ \\citep{Weymann_etal.1991}.  The absorption troughs are broader than 2000~km~s$^{-1}$, sometimes as broad as $\\sim0.1c$, and are presumably due to the intervening components of the outflow originating in the activity of central engines.  The fact that the most luminous quasars showing BALs more frequently \\citep{Ganguly_etal.2007} is consistent with the strong radiation-pressure-driven outflows suggested by simulation-based studies for accretion phenomena (e.g., \\cite{Proga_etal.2000,Ohsuga2007}).  The observed maximum velocity of absorption as function of luminosity has an upper envelope, which can be interpreted as the terminal velocity of radiative driven wind \\citep{Ganguly_etal.2007,Laor&Brandt2002}.  The intrinsic percentage of quasars with BALs is $\\sim$20\\% (e.g., \\cite{Hewett&Foltz2003}).  This percentage means that the BAL phenomenon takes one of major roles in quasars'~activities.  However, the principal parameter determining the finding of BAL features is still unknown.  In the most widely accepted scenario, this percentage represents the covering factor of an outflowing BAL wind, which is preferentially equatorial, and the BAL features can be observed when the accretion disk is almost edge-on to the line of sight, based on spectropolarimetric measurements \\citep{Goodrich&Miller1995,Cohen_etal.1995} and a theoretical disk wind model \\citep{Murray_etal.1995}.  However, some radio observations provided a counterargument to this paradigm.  \\citet{Zhou_etal.2006} and \\citet{Ghosh&Punsly2007} found several BAL quasars with rapid radio variability that indicated very high brightness temperatures, which require Doppler beaming on jets with inclinations of less than \\timeform{35D}, i.e., a nearly face-on view of the accretion disk.  \\citet{Becker_etal.2000} found that about one-third of the radio-detected BAL quasars showed flat radio spectra ($\\alpha>-0.5$, $S_\\nu \\propto \\nu^{\\alpha}$), preferring pole-on jets because this geometry tends to make radio sources core-dominated by significant Doppler effect only on nuclear jets.  Also in optical spectropolarimetry, \\citet{Brotherton_etal.2006} found an electric vector nearly parallel to a large-scale jet axis, implying that BAL outflow is not equatorial, in a Fanaroff-Riley Class II~(FR~II) radio galaxy as a BAL quasar.  Thus, a pole-on outflow would be necessary for at least some of the known radio-emitting BAL quasars.  These results support an alternative proposal that BALs are not closely related to inclinations, and may be associated with a relatively short-lived (possibly episodic) evolutionary phase with a large BAL wind-covering fraction (e.g., \\cite{Briggs_etal.1984,Gregg_etal.2000}).  \\citet{Gregg_etal.2006} pointed out the rarity of FR~II/BAL quasars and their observed anticorrelation between the balnicity index and radio loudness, and suggested that these properties can be explained naturally by an evolutionary scheme.  \\citet{Montenegro-Montes_etal.2008} indicated that many radio-emitting BAL quasars share several radio properties common to young radio sources like Compact Steep Spectrum~(CSS) or Gigahertz-Peaked Spectrum~(GPS) sources.  Thus, `inclination angle' and `evolutionary phase' are two of the most important aspects in understanding BAL quasars.  Very-long-baseline interferometry~(VLBI) instruments in radio wavelengths provide exclusive and crucial opportunities to obtain spatial information about AGNs at milli-arcsecond~(mas) scales by direct measurement, which should also be useful for investigating BAL quasars.  The inclination can be resolved by determining the viewing angle of the jet axis, which is supposed to be perpendicular to the accretion disk.  Using the framework of Doppler beaming effects, jet axes for many AGNs have been estimated by measurements of jet asymmetry~(advancing speed, brightness, jet length, etc.) by VLBI imaging.  Age as a radio source can be also estimated by measurements of its apparent linear size and expanding speed, or the age of relativistic electrons appearing on synchrotron spectrum (e.g., \\cite{Nagai_etal.2006} and references therein).  VLBI observations for BAL quasars have recently begun to try to study such phenomena.  \\citet{Jiang&Wang2003} observed three BAL quasars with the European VLBI Network~(EVN) at 1.6~GHz and suggested that the jet of J1556+3517 was possibly viewed from nearly pole-on because of a flat spectrum and unresolved core, while that of J1312+2319 may be far from pole-on because of the two-sided structure, and the inclination of J0957+2356 was unclear because of an unresolved steep-spectrum compact source.  \\citet{Liu_etal.2008} observed for eight BAL quasar sample, including both of LoBALs and HiBALs and both of steep- and flat-spectrum sources, with the EVN$+$Multi-Element Radio Linked Interferometer Network~(MERLIN) at 1.6~GHz.  High brightness temperatures and linear polarization in their core components implied a synchrotron origin for the radio emission.  No systematic difference was found in the radio morphology or polarization properties between their limited number of LoBAL/HiBAL or steep/flat-spectrum sources.  \\citet{Kunert-Bajraszewska&Marecki2007} observed 1045+352 with the US Very Long Baseline Array~(VLBA) at 1.7, 5, and 8.4~GHz, and found complicated radio morphology and a projected linear size of only 2.1~kpc, which suggests it might be in an early stage, and consistent with the evolutionary scenario.  These investigations were still inconclusive because the detected jet structures, spatial resolutions, and frequency coverages were inadequate to definitively determine jet properties, and also because of still a small number of objects to conclude the natures of BAL quasars.   In this paper, we report our VLBI observations of 22~BAL quasars at 8.4~GHz by direct measurement at the mas resolution, corresponding to the parsec scale at the distance to these sources.  Our aim is to find BAL quasars that have jets suitable for determining their orientation angles and ages from jet properties for future detailed VLBI imaging studies.  In Section~\\ref{section:sample}, we describe our selection processes, and we present our observations and data reduction procedures in Section~\\ref{section:observationanddatareduction}.  We present the observational results in Section~\\ref{section:result}, and discuss their implications in Section~\\ref{section:discussion}.  In Section~\\ref{section:summary}, we summarize the outcome of our investigation.  Throughout this paper, a flat cosmology is assumed, with $H_0=70$~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_\\mathrm{M}=0.3$, and $\\Omega_\\mathrm{\\Lambda}=0.7$ \\citep{Spergel_etal.2003}.      \\begin{table*} \\centering \\rotatebox[origin=c]{90}{ \\begin{minipage}[c]{1.25\\textwidth} \\caption{Optically-selected, radio-flux-limited BAL quasar sample for OCTAVE observation.}\\label{table1} \\begin{center} \\begin{tabular}{llclccrccc} \\hline\\hline \\multicolumn{1}{c}{SDSS name} & \\multicolumn{1}{c}{FIRST name} & $r_\\mathrm{FIRST}^\\mathrm{SDSS}$ & \\multicolumn{1}{c}{$z$} & $l$ & BAL type & \\multicolumn{1}{c}{$M_i^\\mathrm{SDSS}$} & $I^\\mathrm{FIRST}_\\mathrm{1.4GHz}$ & $S^\\mathrm{FIRST}_\\mathrm{1.4GHz}$ & $\\log{R*}$ \\\\ \\multicolumn{1}{c}{} & \\multicolumn{1}{c}{} & (arcsec) & \\multicolumn{1}{c}{} & (pc mas$^{-1}$) &  & \\multicolumn{1}{c}{(mag)} & (mJy beam$^{-1}$) & (mJy) &  \\\\ \\multicolumn{1}{c}{(1)} & \\multicolumn{1}{c}{(2)} & (3) & \\multicolumn{1}{c}{(4)} & (5) & (6) & \\multicolumn{1}{c}{(7)} & (8) & (9) & (10) \\\\\\hline SDSS J004323.43$-$001552.4 & FIRST J004323.8$-$001548 & 7.6  & 2.798  & 7.9  & H & $-$27.94 & 103  & 115  & 2.7  \\\\ SDSS J021728.62$-$005227.2 & FIRST J021728.6$-$005226 & 0.3  & 2.463  & 8.1  & H & $-$25.98 & 212  & 218  & 3.7  \\\\ SDSS J075628.24$+$371455.6 & FIRST J075628.2$+$371455 & 0.3  & 2.514  & 8.1  & HL? & $-$26.10 & 239  & 247  & 3.7  \\\\ SDSS J080016.09$+$402955.6 & FIRST J080016.0$+$402955 & 0.5  & 2.021  & 8.4  & Hi & $-$26.80 & 190  & 200  & 3.1  \\\\ SDSS J081534.16$+$330528.9 & FIRST J081534.1$+$330529 & 0.5  & 2.426  & 8.1  & nH & $-$27.04 & 328  & 342  & 3.4  \\\\ SDSS J092824.13$+$444604.7 & FIRST J092824.1$+$444604 & 0.0  & 1.904  & 8.4  & nHi & $-$27.16 & 156  & 162  & 2.8  \\\\ SDSS J100515.98$+$480533.2 & FIRST J100515.9$+$480533 & 0.2  & 2.372  & 8.2  & Hi & $-$28.29 & 206  & 209  & 2.7  \\\\ SDSS J101329.92$+$491840.9 & FIRST J101329.9$+$491841 & 0.2  & 2.201  & 8.3  & nHi & $-$26.86 & 267  & 269  & 3.3  \\\\ SDSS J101827.85$+$053030.0 & FIRST J101827.8$+$053029 & 0.2  & 1.938  & 8.4  & Hi & $-$26.62 & 284  & 297  & 3.3  \\\\ SDSS J102027.20$+$432056.2 & FIRST J102027.2$+$432056 & 0.2  & 1.962  & 8.4  & nHi & $-$26.74 & 108  & 110  & 2.9  \\\\ SDSS J103038.38$+$085324.9 & FIRST J103038.3$+$085324 & 0.6  & 1.750  & 8.5  & nHi & $-$26.57 & 108  & 172  & 3.0  \\\\ SDSS J104257.58$+$074850.5 & FIRST J104257.5$+$074850 & 0.3  & 2.665  & 8.0  & H & $-$27.32 & 374  & 382  & 3.5  \\\\ SDSS J105726.62$+$032448.0 & FIRST J105726.6$+$032448 & 0.5  & 2.832  & 7.3  & H & $-$26.90 & 138  & 157  & 3.2  \\\\ SDSS J110344.53$+$023209.9 & FIRST J110344.5$+$023209 & 0.2  & 2.514  & 8.1  & H & $-$27.63 & 163  & 166  & 2.9  \\\\ SDSS J111914.32$+$600457.2 & FIRST J111914.3$+$600457 & 0.1  & 2.646  & 8.0  & nH & $-$28.97 & 186  & 192  & 2.5  \\\\ SDSS J115944.82$+$011206.9 & FIRST J115944.8$+$011206 & 0.2  & 2.000  & 8.4  & Hi & $-$28.40 & 267  & 268  & 2.6  \\\\ SDSS J122343.16$+$503753.4 & FIRST J122343.1$+$503753 & 0.1  & 3.488  & 7.3  & H & $-$29.21 & 222  & 229  & 2.7  \\\\ SDSS J122836.92$-$030439.2 & FIRST J122836.9$-$030438 & 0.7  & 1.801  & 8.4  & nHi & $-$26.53 & 144  & 149  & 3.0  \\\\ SDSS J140507.80$+$405657.8 & FIRST J140507.7$+$405658 & 0.3  & 1.993  & 8.4  & Hi & $-$26.31 & 206  & 214  & 3.3  \\\\ SDSS J143243.29$+$410327.9 & FIRST J143243.3$+$410328 & 0.4  & 1.970  & 8.4  & nHi & $-$27.84 & 257  & 262  & 2.8  \\\\ SDSS J151005.88$+$595853.3 & FIRST J151005.4$+$595856 & 4.3  & 1.720  & 8.5  & nHi & $-$27.05 & 182  & 307  & 3.1  \\\\ SDSS J152821.65$+$531030.4 & FIRST J152821.6$+$531030 & 0.4  & 2.822  & 7.8  & H & $-$26.37 & 172  & 183  & 3.6  \\\\\\hline \\end{tabular} \\end{center} \\begin{flushleft} {\\footnotesize Col.~(1) SDSS source name; Col.~(2) Radio counter part from FIRST; Col.~(3) Difference between SDSS--FIRST positions; Col.~(4) Redshift; Col.~(5) Linear scale corresponding to 1~mas; Col.~(6) BAL subtype, listed in \\citet{Trump_etal.2006}.  The code ``n'' denotes a relatively narrow trough; ``HL'' denotes a HiBAL in which broad ($\\geq$1000~km~s$^{-1}$) low-ionization absorption is also seen; ``Hi'' denotes a HiBAL-only object, in which broad low-ionization absorption is not seen even though Mg$_\\mathrm{II}$ is within the spectral coverage; and ``H'' denotes a HiBAL in which the Mg$_\\mathrm{II}$ region is not within the spectral coverage of SDSS or has a very low signal-to-noise ratio and so whether or not the object is a LoBAL as well as a HiBAL is unknown; Col.~(7) $i$-band absolute magnitude listed in \\citet{Trump_etal.2006} and calculated using a flat cosmology of $H_0=70$~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_\\mathrm{M}=0.27$, and $\\Omega_\\mathrm{\\Lambda}=0.73$ \\citep{Spergel_etal.2003}; Col.~(8) 1.4~GHz peak intensity from FIRST data with a $\\sim\\timeform{5\"}$ resolution; Col.~(9) 1.4~GHz flux density from FIRST data; Col.~(10) Radio loudness, the ratio of radio-to-optical flux densities, $R* = f_\\mathrm{5GHz}/f_\\mathrm{2500\\AA}$ (e.g., \\cite{Stocke_etal.1992}), which were calculated from $z$, $S^\\mathrm{FIRST}_\\mathrm{1.4GHz}$, $M_i$, and the assumption of a radio spectral index of $-0.5$ and an optical spectral index of $-0.5$.} \\end{flushleft} \\end{minipage} } \\end{table*} ",
        "conclusions": "\\label{section:discussion} \\subsection{Coexistence of nonthermal jet and BAL outflow}\\label{section:coexistence} We detected 20 radio sources using OCTAVE baselines, which assured brightness temperatures of greater than $10^5$--$10^6$~K.  Although the brightness temperature of a radio supernova at a very early stage could exceed $T_\\mathrm{B}=10^6$~K (e.g., \\cite{Bietenholz_etal.2001}), its expected flux density would be far from sufficient for fringe detection for such distant quasars in our sample (cf.~\\cite{vanDyk_etal.1993}).  The brightness temperatures of stellar components of luminous starbursts are $\\lesssim10^5$~K (cf. \\cite{Lonsdale_etal.1993}), which are less than those of the detected radio counterparts of BAL quasars.  Also, in terms of radio luminosity, even the most radio-luminous starbursts show up to only $\\sim10^{24}$~W~Hz$^{-1}$ (e.g., \\cite{Smith_etal.1998}), in contrast to $\\sim10^{27}$--$10^{28}$~W~Hz$^{-1}$ for our BAL quasar sample.  Therefore, the OCTAVE-detected radio emissions cannot be accounted for by any stellar origin.  We conclude that the OCTAVE-detected radio emissions of the BAL quasars originate in nonthermal jets from AGN activity, as in the cases of other AGN radio sources.  VLBI observations that previously conducted \\citep{Jiang&Wang2003,Kunert-Bajraszewska&Marecki2007,Liu_etal.2008} also support the existence of nonthermal jets in BAL quasars.  The fairly high VLBI-detection rate (20/22) is evidence that BAL outflows, which are inferred from broad troughs in UV spectra, can coexist with nonthermal jets in radio-loud BAL quasars.  This indicates that the accretion disks of BAL quasars can generate radiation-pressure driven strong outflows and magnetic-driven strong jets simultaneously.  It is important to investigate the relationship between the properties of BAL features and radio emissions (cf.~\\cite{Ghosh&Punsly2007}) to understand accretion phenomena in quasars.   \\subsection{Jet properties of observed BAL quasars --- inverted-spectrum sources}\\label{section:jetproperties:invertedspectrum} We found four inverted-spectrum ($\\alpha > 0$) sources (Column~(9) in Table~\\ref{table3}).  Since optically-thin nonthermal synchrotron emission should show steep spectrum of $\\alpha \\leqq -0.5$, the observed inverted spectra should result from a lower-frequency absorption mechanism if a non-simultaneous spectral index is not affected by flux variability.  Absorbed components should dominate the total spectrum for an inverted spectrum to be observed even using the FIRST beam width.  Hence, an inverted spectrum suggests three possibilities: (1)~Doppler beaming effect on jets, (2)~a young~(compact) radio source, or (3)~artificially made due to flux variability, as follows.  Doppler boosting can apparently enhance only optically-thick nuclear components beyond extended jets that would have been decelerated and optically-thin: the nuclear components tend to be synchrotron self-absorbed because of high-brightness temperatures, and show inverted spectrum at lower frequencies.  An adequate Doppler beaming effect requires jets that are nearly aligned with our line-of-sight; this implies a face-on viewed accretion disk, which is inconsistent with the widely accepted scheme of equatorial BAL outflows.  The presence of Doppler beaming effect has already been inferred from rapid flux variation between the NRAO VLA Sky Survey (NVSS; \\cite{Condon_etal.1998}) and the FIRST in several BAL quasars \\citep{Zhou_etal.2006,Ghosh&Punsly2007}.  It is important to confirm the Doppler beaming in also by VLBI imaging in a future.  Young radio sources also can make their spectra nuclear-dominated, because extended jets would not yet been developed.  Radio sources of $\\sim100$--1000~pc or less could show peaked spectra, resulting inverted spectra possibly at $\\sim1$~GHz; the spectral-peak frequency is roughly determined by the linear size of radio structure (e.g., \\cite{Snellen_etal.2000}).  The explanations of the relation between linear size and spectral-peak frequency have been suggested using synchrotron self-absorption~\\citep{Odea&Baum1997} and free--free absorption~\\citep{Bicknell_etal.1997}.  \\citet{Montenegro-Montes_etal.2008} presented many radio sources of BAL quasars showing such peaked spectra.  The ages could be determined from the linear size and expanding speed of radio structures; it is important to measure the expanding speed by VLBI-imaging monitor in a future.  For example, in a typical VLBI scale, an apparent expanding rate of 0.1~mas yr$^{-1}$ and a linear size of two-sided structure of 100~pc leads to an estimated age of 500~yr.  This situation can be adapted to the evolutionary scenario of BAL quasars in a relatively short-lived phase (e.g., \\cite{Briggs_etal.1984,Gregg_etal.2000}, and see also \\cite{Gregg_etal.2006}).    Inverted spectra artificially-made due to flux variability between the observations of the OCTAVE at 8.4~GHz and the VLA at 1.4~GHz ($\\sim10$~yr) cannot be ruled out.  A fraction of BAL quasars are significantly variable in the radio band \\citep{Zhou_etal.2006,Ghosh&Punsly2007}.  A flux variability of $>$10\\% in 10~yr at 8.4~GHz for a 100-mJy source at $z=2$ implies $T_\\mathrm{B} > 10^{11.3}$~K, which would be above the inverse-Compton limit \\citep{Readhead1994} and require Doppler beaming effect.  A flux variability of 10\\% corresponds to a change of spectral index of only $\\sim0.05$.  That means that even if the observed spectra had been artificially made, we have the same conclusion that the four inverted-spectrum sources are possibly Doppler-boosted.  {\\it As a result, we conclude that the inverted-spectrum sources are interpreted as Doppler-beamed, pole-on-viewed relativistic jet sources or young radio sources such as CSSs and GPSs: single edge-on geometry cannot descibe all BAL quasars.}  We also check flux variation between NVSS and FIRST at 1.4~GHz for our BAL quasar sample.  Using the same manner as \\citet{Ghosh&Punsly2007}, we found two sources, FIRST J075628.2$+$371455 and FIRST J122836.9$-$030438, with significant flux variation of $4.8\\sigma$ and $3.2\\sigma$; brightness temperatures were derived to be $10^{15.3}$~K and $10^{13.2}$~K, respectively.  Both of the radio sources were very compact (2~mas or less) in OCTAVE observations, and the latter showed an inverted spectral index ($\\alpha=+2.0$).  These situations are consistent with the presence of highly Doppler boosting on pole-on jets.  It is important to confirm them also by VLBI imaging in a future.     \\subsection{Jet properties of observed BAL quasars --- steep-spectrum sources}\\label{section:jetproperties:steepspectrum} The majority of the detected BAL quasars showed steep ($\\alpha<0$) radio spectra.  It was unclear whether the weaker radio flux densities were observed at 8.4~GHz because extended structures were resolved out, or were due to intrinsically steep spectra on compact components.  Hence, we do not discuss further in detail for the OCTAVE results.  We have already stressed the importance of VLBI imaging observations for inverted-spectrum sources, and steep-spectrum sources are also needed to be VLBI-observed because they are the majority of BAL quasars and are also compact in most cases \\citep{Becker_etal.2000}.  The correlated flux densities in the OCTAVE baselines were not much smaller than the $10^{-1}$-times the VLA peak intensities in most sources, despite the fact that the beam area of the OCTAVE was $\\sim10^{-5}$-times that of VLA.  This indicates that radio-emitting origins considerably concentrated in parsec-scale components, which should be investigated using VLBI.  As the evolutionary scenario, BAL quasars might associate young radio sources \\citep{Briggs_etal.1984,Gregg_etal.2000,Gregg_etal.2006}, which should be optically-thin small radio lobes seen in compact steep spectrum~(CSS; \\cite{Odea1998} for a review) objects.  Many CSS sources have been revealed as young ($<10^5$~yr) radio galaxies by VLBI-imaging studies (e.g., \\cite{Murgia2003,Nagai_etal.2006}).  VLBI-imaging studies may provide evidence supporting the evolutionary scenario for BAL quasars, if the steep-spectrum radio sources of BAL quasars are CSS objects \\citep{Kunert-Bajraszewska&Marecki2007}.       \\subsection{Implications of OCTAVE observations}\\label{section:implications} Our OCTAVE observations have many implications for the study of BAL quasars.  VLBI images of only four BAL quasars have been published before the OCTAVE observations \\citep{Jiang&Wang2003,Kunert-Bajraszewska&Marecki2007}.  Our OCTAVE observations have dramatically increased the number of VLBI-detected BAL quasars, and have established that this AGN subclass includes a non-trivial number of radio-loud objects that can be directly imaged at mas resolutions (corresponding to parsec scales) as well as objects in the other AGN classes.  The orientation angle and the age as radio sources of BAL quasars should be determined from jet properties in subsequent detailed VLBI imaging studies.  The OCTAVE strongly recommends that multi-epoch and multi-frequency (and polarimetric) VLBI observations should be performed for these sources.  Multi-epoch imaging can measure the proper motion of jets or lobes to estimate the inclination or kinematic age.  Multi-frequency imaging can measure the spectral-index profiles along approaching and receding jets or lobes, which can discriminate between Doppler-boosted self-absorbed jets and free--free absorbed jets obscured by thermal plasma.  BAL outflows could be free--free absorber along nonthermal jets; the projected one-dimensional profile of thermal BAL outflows could be studied in mas resolutions using VLBIs at multi-frequency at by measuring opacity gradient along the jets.  The free--free absorber also could occur Faraday rotation to the vector of polarization axis toward jets, which is an another powerful tool to investigate the spatial profile of thermal BAL outflows.  In multi-frequency VLBI observations, such processes should offer exclusive probes to the parsec-scale profile of BAL outflows.  On the basis of the results of the OCTAVE observations, we are observing some of the OCTAVE sample by multi-frequency VLBI imaging.  It will be important to compare the results of inclination measurements based on the VLBIs and optical spectropolarimetry, and to investigate the relation among UV/optical spectra, pc-scale geometry, and ages.       Two of the most important questions in understanding BAL quasars are: (1)~Are they viewed at nearly edge-on? (2) Are they in a short-lived evolutionary phase?  Our OCTAVE observations have increased the number of VLBI-detected BAL quasars, which offers more chance to determine their orientations and ages by subsequent VLBI-imaging studies to conclude the two pictures.  We detected 20 out of the radio-brightest 22~sources selected from the counterparts of SDSS BAL quasars.  We concluded that nonthermal pc-scale jets and thermal BAL outflows can coexist in these radio-loud BAL quasars simultaneously.  BAL quasars have become targets of VLBIs to be revealed in pc scale by direct imaging, as in the cases of other AGN classes.  We also found four inverted-spectrum sources, which are interpreted as Doppler-beamed, pole-on-viewed relativistic jet sources or young radio sources: single edge-on geometry cannot describe all BAL quasars.     \\bigskip  We thank M.~S.~Brotherton for very useful comments that improved the presentation of the paper.  We also thank K.~Asada for very useful comments.  This work was partially supported by a Grant-in-Aid for Young Scientists~(B; 18740107, A.~D.), a Grant-in-Aid for Scientific Research~(C; 21540250, A.~D.) and a Grant-in-Aid for Scientific Research~(C; 20540233, K.~W.) from the Japanese Ministry of Education, Culture, Sports, Science, and Technology~(MEXT).  We are grateful to all the staffs and students involved in the development and operation of the OCTAVE.  The OCTAVE project has been developed as a subproject under the VLBI Exploration of Radio Astrometry~(VERA) project in the National Astronomical Observatory of Japan~(NAOJ), a branch of the National Institutes of Natural Sciences~(NINS).  The optical-fiber networks for OCTAVE have been provided by the GALAXY project\\footnote{The GALAXY lines connecting Usuda 64~m and Nobeyama 45~m to correlators had been closed since April 2008.} supported by NTT Corporation \\citep{Uose_etal.2002}, the Japan Gigabit Network-2 (JGN2) project operated by the National Institute of Information and Communications Technology~(NICT), and the Science Information NETwork~3~(SINET3) operated by the National Institute of Informatics~(NII).  The OCTAVE array consists of six contributed antennas: the Usuda 64~m of the Japan Aerospace Exploration Agency~(JAXA), the Kashima 34~m of the NICT, the Nobeyama 45~m of the Nobeyama Radio Observatory~(NRO) in the NAOJ, the Tsukuba 32~m of the Geographical Survey Institute~(GSI), the Yamaguchi 32~m (operated by Yamaguchi University) of the NAOJ, and the Gifu 11~m of Gifu University.  Time allocation and array operation for OCTAVE are being carried out under the framework of the Japanese VLBI Network~(JVN) project, which is led by the NAOJ, Hokkaido University, Gifu University, Yamaguchi University, and Kagoshima University, in cooperation with the Institute of Space and Astronautical Science~(ISAS)/JAXA, GSI and NICT.  We used the US National Aeronautics and Space Administration's (NASA's) Astrophysics Data System~(ADS) abstract service, the NASA/IPAC Extragalactic Database (NED), which is operated by the Jet Propulsion Laboratory~(JPL).  In addition, we used the Astronomical Image Processing System~(AIPS) software developed at the US  National Radio Astronomy Observatory~(NRAO), a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc."
    },
    "0910/0910.5133_arXiv.txt": {
        "abstract": "The Galactic bulge  is the central spheroid of  our Galaxy, containing about one quarter of the total  stellar mass of the Milky Way (M$_{\\rm bulge}=1.8   \\times  10^{10}   M_\\odot$;  Sofue,   Honma   \\&  Omodaka 2009). Being older than the disk, it is the first massive component of the Galaxy to have collapsed into stars.  Understanding its structure, and the  properties of its  stellar population, is therefore  of great relevance  for galaxy  formation models.   I will  review  our current knowledge of  the bulge properties, with special  emphasis on chemical abundances,   recently    measured   for   several    hundred   stars. ",
        "introduction": "The near infrared images from the COBE/DIRBE experiment clearly showed that our Galaxy has a boxy  shaped bulge (Dwek et al.  1995). Isophote deprojection revealed  a barlike nature, with  axes ratios 1:0.33:0.23 and  with  the  near side  on  the  $1^{\\rm  st}$ quadrant,  at  $\\sim 20^\\circ$ from the Sun-Galactic center direction.  These findings were later confirmed  by several authors (e.g., Babusiaux  \\& Gilmore 2005; Rattenbury et  al.  2007a, and references therein),  hence the prolate nature of  the bulge is now  widely accepted. The scale  length of the {\\it main} bar is $\\sim 1.5$ kpc.  A smaller bar (scale $\\sim 600$ pc) seems to be also present in  the inner bulge (Alard 2001, Nishiyama et al.   2005),  though  further   studies  are  needed  to  confirm  and characterize this structure.  A  striking feature recently  discovered in  the outer  bulge suggests that the Galactic  bulge is all but a  simple prolate spheroid.  Along the minor axis, at distances in  excess of $\\sim 700$ pc, a double red clump  is  clearly  visible  in  several  independent  sets  of  data, symmetric  both  at  positive  and negative  latitudes  (McWilliam  \\& Zoccali  2009). The double  clump disappears  outside the  minor axis, leaving only the  brighter of the two at  positive longitudes, and the fainter of  the two at negative longitudes.   Photometric data mapping the  whole bulge  area are  presently  available only  from the  2MASS survey,  which  is  not  faint  enough  to reach  the  red  clump  for $|b|<3^\\circ$.  The  {\\it Vista Variable  in the Via  L\\'actea} survey (Minniti et al. 2009) will solve this problem, mapping the whole bulge to much fainter magnitudes ($K_s  \\sim 20$ in the coadded images), and allowing  to deproject  its  3D structure  by  means of  its RR  Lyrae variable stars. ",
        "conclusions": "The nature  of the  Galactic bulge is  somehow puzzling.   Its stellar population is old ($10-12$ Gyr)  and it has a metallicity distribution compatible  with  chemical  enrichment  models assuming  a  fast  star formation.   The abundance  ratio  of alpha  elements  over iron  also supports a  short star formation timescale, certainly  more rapid than that of the thin disk, and  possibly more rapid than that of the thick disk. The presence of a  radial metallicity gradient, at least outside $\\sim 600$ pc, favors a formation scenario via dissipational collapse, rather than secular evolution of the disk. Nevertheless, the bulge has the  shape of a  bar (perhaps  including some  X-shape feature  in the outer part) and a  cylindrical rotation velocity, both characteristics of  a  {\\it  pseudobulge}  formed  via dynamical  heating  of  a  bar, resulting from disk secular evolution.  A possible  solution to  these conflicting results  may come  from the confirmation  of a  double component  bulge, as  already  suggested by several studies (e.g., Soto et  al. 2007, Hill et al.  2009, Babusiaux et al. 2009) and seen in several bulges of external galaxies (Peletier et al. 2007).  In any case,  it is now clear from  several independent evidences that the Galactic  bulge is  a complex structure.   There is  a metallicity gradient in the outer region that seems not to be present in the inner region.   There are  indications that  the  MDF in  Baade's Window  is bimodal, with  each of the two component  having different kinematics. Stellar  ages are  predicted to  be  different along  the minor  axis, compared to the edges of the  bar. Even the morphology itself does not seem to be simply that of a bar, but rather something like an X-shape. The properties of  the stellar population in Baade's  Window cannot be considered  as   representative  of  the  whole   bulge:  larger  area photometric  and  spectroscopic {\\it  maps}  are  needed  in order  to understand the  bulge structure and  origin.  The VVV survey,  and its spectroscopic followups, will certainly reserve many surprises in this sense."
    },
    "0910/0910.0974.txt": {
        "abstract": "Due to the interaction of physics and astrophysics we are witnessing in these years a splendid synthesis of theoretical, experimental and observational results originating from three fundamental physical processes. They were originally proposed by Dirac, by Breit and Wheeler and by Sauter, Heisenberg, Euler and Schwinger. For almost seventy years they have all three been followed by a continued effort of experimental verification on Earth-based experiments. The Dirac process, $e^+e^-\\rightarrow2\\gamma$, has been by far the most successful. It has obtained extremely accurate experimental verification and has led as well to an enormous number of new physics in possibly one of the most fruitful experimental avenues by introduction of storage rings in Frascati and followed by the largest accelerators worldwide: DESY, SLAC etc. The Breit--Wheeler process, $2\\gamma\\rightarrow e^+e^-$, although conceptually simple, being the inverse process of the Dirac one, has been by far one of the most difficult to be verified experimentally. Only recently, through the technology based on free electron X-ray laser and its numerous applications in Earth-based experiments, some first indications of its possible verification have been reached. The vacuum polarization process in strong electromagnetic field, pioneered by Sauter, Heisenberg, Euler and Schwinger, introduced the concept of critical electric field $E_c=m_e^2c^3/(e\\hbar)$. It has been searched without success for more than forty years by heavy-ion collisions in many of the leading particle accelerators worldwide.  The novel situation today is that these same processes can be studied on a much more grandiose scale during the gravitational collapse leading to the formation of a black hole being observed in Gamma Ray Bursts (GRBs). This report is dedicated  to the scientific race. The theoretical and experimental work developed in Earth-based laboratories is confronted with the theoretical interpretation of space-based observations of phenomena originating on cosmological scales. What has become clear in the last ten years is that all the three above mentioned processes, duly extended in the general relativistic framework, are necessary for the understanding of the physics of the gravitational collapse to a black hole. Vice versa, the natural arena where these processes can be observed in mutual interaction and on an unprecedented scale, is indeed the realm of relativistic astrophysics.  We systematically analyze the conceptual developments which have followed the basic work of Dirac and Breit--Wheeler. We also recall how the seminal work of Born and Infeld inspired the work by Sauter, Heisenberg and Euler on effective Lagrangian leading to the estimate of the rate for the process of electron--positron production in a constant electric field. In addition of reviewing the intuitive semi-classical treatment of quantum mechanical tunneling for describing the process of electron--positron production, we recall the calculations in \\emph{Quantum Electro-Dynamics} of the Schwinger rate and effective Lagrangian for constant electromagnetic fields. We also review the electron--positron production in both time-alternating electromagnetic fields, studied by Brezin, Itzykson, Popov, Nikishov and Narozhny, and the corresponding processes relevant for pair production at the focus of coherent laser beams as well as electron beam--laser collision. We finally report some current developments based on the general JWKB approach which allows to compute the Schwinger rate in spatially varying and time varying electromagnetic fields.  We also recall the pioneering work of Landau and Lifshitz, and Racah on the collision of charged particles as well as experimental success of AdA and ADONE in the production of electron--positron pairs.  We then turn to the possible experimental verification of these phenomena. We review: (A) the experimental verification of the $e^+e^-\\rightarrow 2\\gamma$ process studied by Dirac. We also briefly recall the very successful experiments of $e^+e^-$ annihilation to hadronic channels, in addition to the Dirac electromagnetic channel; (B) ongoing Earth based experiments to detect electron--positron production in strong fields by focusing coherent laser beams and by electron beam--laser collisions; and (C) the multiyear attempts to detect electron--positron production in Coulomb fields for a large atomic number $Z>137$ in heavy ion collisions. These attempts follow the classical theoretical work of Popov and Zeldovich, and Greiner and their schools.  We then turn to astrophysics. We first review the basic work on the energetics and electrodynamical properties of an electromagnetic black hole and the application of the Schwinger formula around Kerr--Newman black holes as pioneered by Damour and Ruffini. We only focus on black hole masses larger than the critical mass of neutron stars, for convenience assumed to coincide with the Rhoades and Ruffini upper limit of 3.2 $M_\\odot$. In this case the electron Compton wavelength is much smaller than the spacetime curvature and all previous results invariantly expressed can be applied following well established rules of the equivalence principle. We derive the corresponding rate of electron--positron pair production and introduce the concept of dyadosphere. We review recent progress in describing the evolution of optically thick electron--positron plasma in presence of supercritical electric field, which is relevant both in astrophysics as well as ongoing laser beam experiments. In particular we review recent progress based on the Vlasov-Boltzmann-Maxwell equations to study the feedback of the created electron--positron pairs on the original constant electric field. We evidence the existence of plasma oscillations and its interaction with photons leading to energy and number equipartition of photons, electrons and positrons. We finally review the recent progress obtained by using the Boltzmann equations to study the evolution of an electron--positron-photon plasma towards thermal equilibrium and determination of its characteristic timescales. The crucial difference introduced by the correct evaluation of the role of two and three body collisions, direct and inverse, is especially evidenced. We then present some general conclusions.  The results reviewed in this report are going to be submitted to decisive tests in the forthcoming years both in physics and astrophysics. To mention only a few of the fundamental steps in testing in physics we recall the starting of experimental facilities at the National Ignition Facility at the Lawrence Livermore National Laboratory as well as corresponding French Laser the Mega Joule project. In astrophysics these results will be tested in galactic and extragalactic black holes observed in binary X-ray sources, active galactic nuclei, microquasars and in the process of gravitational collapse to a neutron star and also of two neutron stars to a black hole giving origin to GRBs. The astrophysical description of the stellar precursors and the initial physical conditions leading to a gravitational collapse process will be the subject of a forthcoming report. As of today no theoretical description has yet been found to explain either the emission of the remnant for supernova or the formation of a charged black hole for GRBs. Important current progress toward the understanding of such phenomena as well as of the electrodynamical structure of neutron stars, the supernova explosion and the theories of GRBs will be discussed in the above mentioned forthcoming report. What is important to recall at this stage is only that both the supernovae and GRBs processes are among the most energetic and transient phenomena ever observed in the Universe: a supernova can reach energy of $\\sim 10^{54}$ ergs on a time scale of a few months and GRBs can have emission of up to $\\sim 10^{54}$ ergs in a time scale as short as of a few seconds. The central role of neutron stars in the description of supernovae, as well as of black holes and the electron--positron plasma, in the description of GRBs, pioneered by one of us (RR) in 1975, are widely recognized. Only the theoretical basis to address these topics are discussed in the present report. ",
        "introduction": "\\label{qedvac}  \\emph{Quantum Electro-Dynamics} (QED), the quantum theory of electrons, positrons, and photons, was established by by Tomonaga \\cite{1946PThPh...1...27T}, Feynman \\cite{1948RvMP...20..367F,1949PhRv...76..749F,1949PhRv...76..769F}, Schwinger \\cite{1948PhRv...74.1439S,1949PhRv...75..651S,1949PhRv...76..790S} and Dyson \\cite{1949PhRv...75..486D,1949PhRv...75.1736D} and others in the 1940's and 1950's \\cite{Schwinger1958}. For decades, both theoretical computations and experimental tests have been developed to great perfection. It is now one of the fundamental pillars of the theory of the microscopic world. Many excellent monographs have been written \\cite{1959ittq.book.....B,Bjorken1998,Bjorken1965,Feynman1998,Feynman1965,1982els..book.....B,Itzykson2006,Lee1990,1951PhRv...82..664S,1954PhRv...93..615S,1954PhRv...94.1362S,Schwinger1970,Schwinger1998,1995qtf..book.....W,Kleinert1990}% , so the concepts of the theory and the techniques of calculation are well explained. On the basis of this material, we review some aspects and properties of the QED that are relevant to the subject of the present review.  QED combines a relativistic extension of quantum mechanics with a quantized electromagnetic field. The nonrelativistic system has a unique ground state, which is the state with no particle, the \\emph{vacuum state\\/}. The excited states contain a fixed number of electrons and an arbitrary number of photons. As electrons are allowed to become relativistic, their number becomes also arbitrary, and it is possible to create pairs of electrons and positrons.  In the modern functional integral description, the nonrelativistic system is described by a given set of fluctuating particle orbits running forward in time. If the theory is continued to an imaginary time, in which case one speaks of a \\emph{Euclidean formulation\\/}, the nonrelativistic system corresponds to a canonical statistical ensemble of trajectories.  In the relativistic system, the orbits form worldlines in four-dimensional space-time which may run in any time direction, in particular they may run backwards in time, in which case the backward parts of a line correspond to positrons. The number of lines is arbitrary and the Euclidean formulation corresponds to a grand-canonical ensemble. The most efficient way of describing such an ensemble is by a single fluctuating field \\cite{Kleinert1990}.  The vacuum state contains no physical particles. It does, however, harbor zero-point oscillations of the electron and photon fields. In the worldline description, the vacuum is represented by a grand-canonical ensemble of interacting closed world lines. These are called \\emph{virtual particles\\/}. Thus the vacuum contains the full complexity of a many-body problem so that one may rightfully say that \\emph{the vacuum is the world\\/} \\cite{Streater2000}. In the Fourier decomposition of the fluctuating fields, virtual particles correspond to Fourier components, or \\emph{modes\\/}, in which the 4-vectors of energy and momentum $k^{\\mu}\\equiv(k^{0},\\mathbf{k})\\equiv({\\mathcal{E}% },\\mathbf{k})$ do not satisfy the mass-shell relation \\begin{equation} k^{2}\\equiv(k^{0})^{2}-c^{2}|\\mathbf{k}|^{2}={\\mathcal{E}}^{2}-c^{2}% |\\mathbf{k}|^{2}=m_{e}^{2}c^{4}, \\label{massshellrelation}% \\end{equation} valid for real particles.  The only way to evaluate physical consequences from QED is based on the smallness of the electromagnetic interaction. It is characterized by the dimensionless fine structure constant $\\alpha$. All theoretical results derived from QED are found in the form of series expansions in powers of $\\alpha$, which are expansions around the non-interacting system. Unfortunately, all these expansions are badly divergent (see e.g. Section 4.62 in \\cite{Kleinert2008}). The number of terms contributing to the same order of $\\alpha$ grows factorially fast, i.e., faster than any exponential, leading to a zero radius of convergence. Fortunately, however, the coupling $\\alpha$ is so small that the series possess an apparent convergence up to order $1/\\alpha\\approx137$, which is much higher than will be calculable for a long time to come (see e.g. Section 4.62 in \\cite{Kleinert2008}). With this rather academic limitation, perturbation expansions are well defined.  In perturbation expansions, all physical processes are expressible in terms of \\emph{Feynman diagrams\\/}. These are graphic representations of the interacting world lines of all particles. Among these lines, there are some which run to infinity. They satisfy the mass shell relation (\\ref{massshellrelation}) and describe real particles observable in the laboratory. Those which remain inside a finite space-time region are virtual.  The presence of virtual particles in the perturbation expansions leads to observable effects. Some of these have been measured and calculated with great accuracy. The most famous examples are  \\begin{enumerate} \\item the electrostatic polarizability of quantum fluctuations of the QED vacuum has been measured in the Lamb shift \\cite{1947PhRv...72..241L,1953PhRv...89...98T}.  \\item the anomalous magnetic moment of the electron \\cite{1948PhRv...73..416S,1949PhRv...76..769F,1972PhR.....3..193L,1974PhRvD..10.4007C,1977PhLB...68..191B}.  \\item the dependence of the electric charge on the distance. It is observed by measuring cross-sections of electron--positron collisions, most recently in the L3-experiments at the \\emph{Large Electron-Positron Collider\\/} (LEP) at CERN \\cite{L3CERN}.  \\item the \\emph{Casimir effect\\/} caused by virtual photons, i.e., by the fluctuations of the electromagnetic field in the QED vacuum \\cite{Casimir1948,Fierz1960}. It causes an attractive force \\cite{1958Phy....24..751S,1997PhRvL..78....5L,1998PhRvL..81.4549M} between two uncharged conducting plates in the vacuum (see also \\cite{2001PhRvE..63e1101H,2000PhRvA..62e2109H,2004PhRvA..69b2117C,Xue:1987fa,Xue:1988xj,1986AcPSn..34.1084X,Zheng1993}). \\end{enumerate}  There are, of course, many other discussions of the effects of virtual particles caused either by external boundary conditions or by external classical fields \\cite{1992PNAS...89.4091S,1992PNAS...8911118S,1993PNAS...90..958S,1993PNAS...90.2105S,1993PNAS...90.4505S,1993PNAS...90.7285S,1994PNAS...91.6473S,Bordag1999,Bordag1996,2004JPhA...37R.209M,Milton2001,2001PhLB..508..211X,2003PhRvD..68a3004X,2003MPLA...18.1325X,2005GReGr..37..857X}.  An interesting aspect of virtual particles both theoretically and experimentally is the possibility that they can become real by the effect of external fields. In this case, real particles are excited out of the vacuum. In the previous Section~\\ref{sauter} and \\ref{hew-effecitve}, we have shown that this possibility was first pointed out in the framework of quantum mechanics by Klein, Sauter, Euler and Heisenberg \\cite{1929ZPhy...53..157K,1931ZPhy...73..547S,1931ZPhy...69..742S,1936ZPhy...98..714H} who studied the behavior of the Dirac vacuum in a strong external electric field. Afterward, Schwinger studied this process and derived the probability (\\textit{Schwinger formula}) in the field theory of Quantum Electro-Dynamics, which will be described in this chapter. If the field is sufficiently strong, the energy of the vacuum can be lowered by creating an electron--positron pair. This makes the vacuum unstable. This is the \\emph{Sauter-Euler-Heisenberg-Schwinger process\\/} for electron--positron pair production. There are many reasons for the interest in the phenomenon of pair production in a strong electric field. The most compelling one is that now both laboratory conditions and astrophysical events provide possibilities for observing this process.  In the following chapters, in addition to reviewing the \\textit{Schwinger formula} and QED-effective Lagrangian in constant electromagnetic fields, we will also derive the probability of pair production in an alternating field, and discuss theoretical studies of pair production in (i) electron-beam--laser collisions and (ii) superstrong Coulomb potential. In addition, the plasma oscillations of electron--positron pairs in electric fields will be reviewed in Section \\ref{time-independent}. The rest part of this chapter, we shall use natural units $\\hbar=c=1$.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%   \\subsection{Basic processes in Quantum Electro-Dynamics}\\label{qedprocesse}  The total Lagrangian describing the interacting system of photons, electrons, and positrons reads, see e.g. \\cite{1982els..book.....B} \\begin{equation} {\\mathcal{L}}={\\mathcal{L}}_{0}^{\\gamma}+{\\mathcal{L}}_{0}^{e^{+}e^{-}% }+{\\mathcal{L}}_{\\mathrm{int}}, \\label{qcdl}% \\end{equation} where the free Lagrangians ${\\mathcal{L}}_{0}^{e^{+}e^{-}}$ and ${\\mathcal{L}% }_{0}^{\\gamma}$ for electrons and photons are expressed in terms of quantized Dirac field $\\psi(x)$ and quantized electromagnetic field $A_{\\mu}(x)$ as follows: \\begin{align} {\\mathcal{L}}_{0}^{e^{+}e^{-}}  &  =~\\,\\bar{\\psi}(x)(i\\gamma^{\\mu}% \\partial_{\\mu}-m_{e})\\psi(x),\\label{L0pair}\\\\ {\\mathcal{L}}_{0}^{\\gamma}~~  &  =-{\\frac{1}{4}}F_{\\mu\\nu}(x)F^{\\mu\\nu }(x)+\\mathrm{gauge\\!-\\!fixing~term}. \\label{L0gamma}% \\end{align} Here $\\gamma^{\\mu}$ are the $4\\times4$ Dirac matrices, $\\bar{\\psi}% (x)\\equiv\\psi^{\\dagger}(x)\\gamma^{0}$, and $F_{\\mu\\nu}=\\partial_{\\mu}A_{\\nu }-\\partial_{\\nu}A_{\\mu}$ denotes the electromagnetic field tensor. Minimal coupling gives rise to the interaction Lagrangian \\begin{equation} {\\mathcal{L}}_{\\mathrm{int}}=-ej^{\\mu}(x)A_{\\mu}(x),\\quad\\quad j^{\\mu}% (x)=\\bar{\\psi}(x)\\gamma^{\\mu}\\psi(x). \\label{L0int}% \\end{equation}  After quantization, the photon field is expanded into plane waves as \\begin{equation} A_\\mu(x)=\\int \\frac{d^3k}{2k_0(2\\pi)^3}\\sum^3_{\\lambda=1}\\left[a^{(\\lambda)}({\\bf k} )\\epsilon_\\mu^{(\\lambda)}({\\bf k} )e^{-ikx} +a^{(\\lambda)\\dagger}({\\bf k} )\\epsilon_\\mu^{(\\lambda)*}({\\bf k} )e^{ikx}\\right], \\label{photonmodes} \\end{equation} where $\\epsilon_\\mu^{(\\lambda)}$ are polarization vectors, and $a^{(\\lambda)}$, $a^{(\\lambda)\\dagger}$ are annihilation and creation operators of photons. The quantized fermion field $\\psi(x)$ has the expansion \\begin{eqnarray} \\!\\!\\!\\!\\!\\!\\!\\!\\!\\psi(x) &=& \\int{\\frac{d^3k}{ (2\\pi)^3}} \\frac{m}{ k^0} \\sum_{\\alpha=1,2}\\Big[b_\\alpha({\\bf k},s_3) u^{(\\alpha)}({\\bf k},s_3)e^{-ikx}+d^\\dagger_\\alpha({\\bf k},s_3)v^{(\\alpha)}({\\bf k},s_3)e^{ikx}\\Big],\\nonumber\\\\ \\label{fermionmodes} \\end{eqnarray} where the four-component spinors $u^{(\\alpha)}({\\bf k},s_3)$, $v^{(\\alpha)}( {\\bf k},s_3)$ are positive and negative energy solutions of the Dirac equation with momentum ${\\bf k}$ and spin component $s_3$. The operators $b({\\bf k},s_3)$, $b^\\dagger({\\bf k},s_3)$ annihilate and create electrons, the operators $d({\\bf k},s_3)$ and $d^\\dagger({\\bf k},s_3)$ do the same for positrons \\cite{1982els..book.....B}.  In the framework of QED the transition probability from an initial to a final state for a given process is represented by the imaginary part of the unitary S-matrix squared \\begin{equation} {\\mathcal P}_{f\\leftarrow i}=\\left\\vert \\left\\langle \\mathrm{{f,out}\\left\\vert Im\\,{\\mathcal S}\\right\\vert {i,in}}\\right\\rangle \\right\\vert ^{2}, \\end{equation} where \\begin{equation} \\mathrm{Im\\,{\\mathcal S}}=(2\\pi)^{4}\\delta^{4}(P_{f}-P_{i})\\left\\vert M_{fi}\\right\\vert , \\end{equation} $M_{fi}$ is called matrix element and $\\delta$-function stays for energy-momentum conservation in the process.  When initial state contains two particles with energies $\\epsilon_{1}$ and $\\epsilon_{2}$, and final state contain arbitrary number of particles having 3-momenta $\\mathbf{p}_{i}^{\\prime}$, the transition probability per unit time and unit volume is given by% \\begin{equation} \\frac{d{\\mathcal P}_{f\\leftarrow i}}{dVdt}=(2\\pi)^{4}\\delta^{4}(P_{f}-P_{i})\\left\\vert M_{fi}\\right\\vert ^{2}\\frac{1}{4\\epsilon_{1}\\epsilon_{2}}% %TCIMACRO{\\dprod \\limits_{i}}% %BeginExpansion {\\displaystyle\\prod\\limits_{i}} %EndExpansion \\frac{d^{3}p_{i}^{\\prime}}{\\left(  2\\pi\\right)  ^{3}2\\epsilon_{i}}. \\label{dif_prob}% \\end{equation} The Lorentz invariant differential cross-section for a given process is then obtained from (\\ref{dif_prob}) by dividing it on the flux density of initial particles% \\begin{equation} d\\sigma=(2\\pi)^{4}\\delta^{4}(P_{f}-P_{i})\\left\\vert M_{fi}\\right\\vert ^{2}\\frac{1}{4I_{kin}}% %TCIMACRO{\\dprod \\limits_{i}}% %BeginExpansionv {\\displaystyle\\prod\\limits_{i}} %EndExpansion \\frac{d^{3}p_{i}^{\\prime}}{\\left(  2\\pi\\right)  ^{3}2\\epsilon_{i}}, \\end{equation} where $p_{1}$ and $p_{2}$\\ are particles' 4-momenta, $m_{1}$ and $m_{2}$\\ are their masses respectively, $I_{kin}=\\sqrt{\\left(  p_{1}p_{2}\\right)  ^{2}-m_{1}% ^{2}m_{2}^{2}}$.  It is useful to work with Mandelstam variables which are kinematic invariants built from particles 4-momenta. Consider the process $A+B\\longrightarrow C+D$. Lorentz invariant variables can be constructed in the following way% \\begin{align} s  &  =\\left(  p_{A}+p_{B}\\right)  ^{2}=\\left(  p_{C}+p_{D}\\right) ^{2},\\nonumber\\\\ t  &  =\\left(  p_{A}+p_{C}\\right)  ^{2}=\\left(  p_{B}+p_{D}\\right)  ^{2},\\\\ u  &  =\\left(  p_{B}+p_{C}\\right)  ^{2}=\\left(  p_{A}+p_{D}\\right) ^{2}.\\nonumber \\end{align} Since any incoming particle can be regarded as outgoing antiparticle, it gives rise to the crossing symmetry property of the scattering amplitude, which is best reflected in the Mandelstam variables. In fact, reactions $A+B\\longrightarrow C+D$, $A+\\bar{C}\\longrightarrow\\bar{B}+D$ or $A+\\bar {D}\\longrightarrow C+\\bar{B}$ where the bar denotes the antiparticle are just different cross-channels of a single general reaction. The meaning of the variables $s,t,u$ changes, but the amplitude is the same.  The S-matrix is computed through the interaction operator as% \\begin{equation} {\\mathcal S}={\\mathcal T}\\exp\\left(  i\\int{\\mathcal{L}}_{\\mathrm{int}}d^{4}x\\right)  , \\end{equation} where ${\\mathcal T}$ is the chronological operator. Perturbation theory is applied, since the fine structure constant is small, while any additional interaction in collision of particles contains the factor $\\alpha$.  A\\ simple and elegant way of computation of the S-matrix and consequently of the matrix element $M_{fi}$\\ is due to Feynman, who discovered a graphical way to depict each QED process, in momentum representation.  In what follows we consider briefly the calculation for the case of Compton scattering process \\cite{1982els..book.....B}, which is given by two Feynman diagrams. Conservation law for 4-momenta is $p+k=p^{\\prime}+k^{\\prime}$, where $p$ and $k$ are 4-momenta of electron and photon respectively, and invariant $I_{kin}=\\frac{1}{4} (s-m_e^{2})^{2}$. After the calculation of traces with gamma-matrices, the final result, expressed in Mandelstam variables, is \\begin{align} |M_{fi}|^{2}  &  =2^{7}\\pi^{2}e^{4}\\left[  \\frac{m_e^{2}}{s-m_e^{2}}+\\frac{m_e^{2}% }{u-m_e^{2}}+\\left( \\frac{m_e^{2}}{s-m_e^{2}}+\\frac{m_e^{2}}{u-m_e^{2}}\\right) ^{2}\\right. \\nonumber\\\\ &  \\left.  -\\frac{1}{4}\\left(  \\frac{s-m_e^{2}}{u-m_e^{2}}+\\frac{u-m_e^{2}}{s-m_e^{2}% }\\right)  \\right]  , \\label{M_fi_gamma1_1}% \\end{align} $s=(p+k)^{2}$, $t=(p-p^{\\prime})^{2}$ and$\\ u=(p-k^{\\prime})^{2}$. Since the differential cross-section is independent of the azimuth of $\\mathbf{p}% _{1}^{\\prime}$ relative to $\\mathbf{p}_{1}$,\\ it is obtained from (\\ref{M_fi_gamma1_1}) as% \\begin{equation} d\\sigma=\\frac{1}{64\\pi}|M_{fi}|^{2}\\frac{dt}{I_{kin}^{2}}. \\label{dsigma} \\end{equation}  In the laboratory frame, where $s-m_e^{2}=2m_ew$, $u-m_e^{2}=-2m_ew^{\\prime}$ and electron is at rest before the collision with photon, the differential cross-section of Compton scattering is thus given by the Klein--Nishina formula \\cite{1929ZPhy...52..853K} \\begin{equation} d\\sigma=\\frac{1}{2}\\left(  \\frac{e^{2}}{m}\\right)  ^{2}\\left(  \\frac {\\omega^{\\prime}}{\\omega}\\right)  ^{2}\\left(  \\frac{\\omega}{\\omega^{\\prime}}+\\frac{\\omega^{\\prime}}% {\\omega}-\\sin^{2}\\vartheta\\right)  , \\label{KN}% \\end{equation} where $\\omega$ and $\\omega^{\\prime}$\\ are frequencies of photon before and after the collision, $\\vartheta$\\ is the angle at which the photon is scattered.  \\subsection{The Dirac and the Breit--Wheeler processes in QED}\\label{DiracBWqed}  We turn now to the formulas obtained within framework of quantum mechanics by Dirac \\cite{1930PCPS...26..361D} and Breit and Wheeler \\cite{1934PhRv...46.1087B} within QED. The crossing symmetry allows to readily write the matrix element for the pair production (\\ref{ee2gamma}) and pair annihilation (\\ref{2gammaee}) processes with the energy-momentum conservation written as $p_++p_-=k_1+k_2$, where $p_+$ and $p_-$ are 4-momenta of the positron and the electron, $k_1$ and $k_2$ are 4-momenta of two photons. It is in fact given by the same formula (\\ref{M_fi_gamma1_1}) with the substitution $p\\rightarrow p_{-},$ $p^{\\prime }\\rightarrow p_{+},$ $k\\rightarrow k_{1},$ $k^{\\prime}\\rightarrow k_{2}$, but with different meaning of the kinematic invariants $s=(p_{-}-k_{1})^{2}$, $t=(p_{-}+p_{+})^{2}$,$\\ u=(p_{-}-k_{2})^{2}$. Matrix elements for Dirac and Breit--Wheeler processes are the same. The differential cross-section of the Dirac process is obtained from (\\ref{M_fi_gamma1_1})\\ with the exchange $s\\leftrightarrow t$ and the invariant $I_{kin}=\\frac{1}{4}t(t-4m_e^{2})$, which leads to (\\ref{Dirac cross-section}). For the case of the Breit--Wheeler process with the invariant $I_{kin}=\\frac{1}{4}t^{2}$, the result is reduced to (\\ref{BW section0}).  Since the Dirac pair annihilation process (\\ref{ee2gamma}) is the inverse of Breit--Wheeler pair production (\\ref{2gammaee}), it is useful to compare the cross-section of the two processes. We note that the squared transition amplitude $|M_{fi}|^2$ must be the same for two processes, due to the CPT invariance. The cross-sections could be different only by kinematics and statistical factors. Let us consider the pair annihilation process in the center of mass system where ${\\mathcal E}={\\mathcal E}_1+{\\mathcal E}_2={\\mathcal E}'_1+{\\mathcal E}'_2$ is the total energy, the initial and final momenta are equal and opposite, ${\\bf p}_1 =-{\\bf p}_2\\equiv {\\bf p}$ and ${\\bf p}'_1 =-{\\bf p}'_2\\equiv {\\bf p}'$. The differential cross-section is given by (\\ref{dsigma}). For the Breit and Wheeler process (\\ref{2gammaee}) of two colliding photons with 4-momenta $k_1$ and $k_2$, the scalar $I^2_{\\gamma\\gamma}=(k_1k_2)^2$. For the Dirac process (\\ref{ee2gamma}) of colliding electron and positron with 4-momenta $p_1$ and $p_2$, the scalar $I^2_{e^+ e^-} = (p_1p_2)^2-m_e^4$. As results, one has \\begin{equation}\\label{I} \\frac{d\\sigma_{\\gamma\\gamma}}{d\\sigma_{e^+ e^-}} =\\frac{I^2_{e^+ e^-}}{I^2_{\\gamma\\gamma}} = \\frac{2(k_1k_2)-4m_e^2}{2(k_1k_2)}= \\frac{{\\mathcal E}^2-2m_e^2}{{\\mathcal E}^2} =\\left(\\frac{|{\\bf p}|}{{\\mathcal E}}\\right)^2=\\hat\\beta^2, \\end{equation} where momenta and energies are related by \\begin{equation} (p_1+p_2)^2=(k_1+k_2)^2=2(k_1k_2)=2{\\mathcal E}^2. \\nonumber \\end{equation} Integrating Eq.~(\\ref{I}) over all scattering angles yields the total cross-section. Whereas the previous $\\sigma_{\\gamma\\gamma}$ required division by a Bose factor 2 for the two identical photons in the final state, the cross-section $\\sigma_{e^+ e^-}$ has no such factor since the final electron and positron are not identical. Hence we obtain \\begin{equation}\\label{sigmaee-sigmagg} \\sigma_{e^+ e^-} = \\frac{1}{2\\hat\\beta^2}\\sigma_{\\gamma\\gamma}. \\end{equation} By re-expressing the kinematic quantities in the laboratory frame, one obtains the Dirac cross-section (\\ref{Dirac cross-section}).  As shown in Eq.~(\\ref{sigmaee-sigmagg}) in the center of mass of the system, the two cross-sections $\\sigma_{e^+ e^-}$ and $\\sigma_{\\gamma\\gamma}$ of the above described phenomena differ only in the kinematics and statistical factor $1/(2\\hat\\beta^{2})$, which is related to the fact that the resulting particles are massless or massive. The process of electron and positron production by the collision of two photons has a kinematic energy threshold, while the process of electron and positron annihilation to two photons has not such kinematic energy threshold. In the limit of high energy neglecting the masses of the electron and positron, $\\hat\\beta\\rightarrow 1$, the difference between two cross-sections $\\sigma_{e^+ e^-}$ and $\\sigma_{\\gamma\\gamma}$ is only the statistical factor $1/2$.  The total cross-sections (\\ref{Dirac section0}) of Breit--Wheeler's and Dirac's process are of the same order of magnitude $\\sim 10^{-25}{\\rm cm}^2$ and have the same energy dependence $1/{\\mathcal E}^2$ above the energy threshold. The energy threshold ($2m_ec^2$) have made until now technically impossible to observe the pair production by the Breit--Wheeler process in laboratory experiments at the intersection of two beams of X-rays. Another reason is of course the smallness of the total cross-section (\\ref{bwsection3}) ($\\sigma_{\\gamma\\gamma}\\lesssim 10^{-25}$cm$^2$) and the experimental limitations on the intensities $I_i$ (\\ref{bwaa}) of the light beams. We shall see however, that this Breit--Wheeler process occurs routinely in the dyadosphere of a black hole. The observations of such phenomena in the astrophysical setting are likely to give the first direct observational test of the validity of the Breit--Wheeler process, see e.g. \\cite{Ruffini2009}.  \\subsection{Double-pair production}\\label{doub}  Following the Breit--Wheeler pioneer work on the process (\\ref{2gammaee}), Cheng and Wu \\cite{1969PhRvL..22..666C,1969PhRvL..23.1311C,1969PhRv..182.1852C,1969PhRv..182.1868C,1969PhRv..182.1873C,1969PhRv..182.1899C} considered the high-energy behavior of scattering amplitudes and cross-section of two photon collision, up to higher order ${\\mathcal O}(\\alpha^4)$ \\cite{1970PhRvD...1.3414C}, \\begin{equation}\\label{gamma4ee} \\gamma_{1}+\\gamma_{2}\\rightarrow e^{+}+e^{-}+ e^{+}+e^{-}. \\end{equation} For this purpose, they calculated the two photon forward scattering amplitude $M_{\\gamma\\gamma}$ (see Eq.~(\\ref{dif_prob})) by taking into account all relevant Feynman diagrams via two electron loops up to the order ${\\mathcal O}(\\alpha^4)$. The total cross-section $\\sigma_{\\gamma\\gamma}$ for photon-photon scattering is related to the photon-photon scattering amplitude $M_{\\gamma\\gamma}$ in the forward direction by the optical theorem, \\begin{equation}\\label{optict} \\sigma_{\\gamma\\gamma}(s)=\\frac{1}{s}{\\rm Im}M_{\\gamma\\gamma}, \\end{equation} where $s$ is the square of the total energy in the center of mass system. They obtained the total cross-section of double pair production (\\ref{gamma4ee}) at high energy $s\\gg 2m_e$, \\begin{equation}\\label{cwresult} \\lim_{s\\rightarrow \\infty}\\sigma_{\\gamma\\gamma}(s)= \\frac{\\alpha^4}{36\\pi m_e^2}[175\\zeta(3)-38]\\sim 6.5 \\mu b, \\end{equation} which is  independent of $s$ as well as helicities of the incoming photons. Up to the $\\alpha^4$ order, Eq.~(\\ref{cwresult}) is the largest term in the total cross-section for photon-photon scattering at very high energy. This can be seen by comparing Eq.~(\\ref{cwresult}) with the cross-section (\\ref{BW section0},\\ref{bwsection3}) of the Breit--Wheeler process (\\ref{2gammaee}), which is the lowest-order process in a photon-photon collision and vanishes as $s\\rightarrow \\infty$. Thus, although the Breit--Wheeler cross-section is of lower order in $\\alpha$, the Cheng-Wu cross-section (\\ref{cwresult}) is larger as the energy becomes sufficiently high. In particular, the Cheng-Wu cross-section (\\ref{cwresult}) exceeds the Breit--Wheeler one (\\ref{BW section0}) as the center of mass energy of the photon ${\\mathcal E}\\ge 0.24$GeV. Note that in the double pair production (\\ref{gamma4ee}), the energy threshold is ${\\mathcal E}\\ge 2m_e$ rather than ${\\mathcal E}\\ge m_e$ in the one pair production of Breit--Wheeler.  In Ref.~\\cite{1970PhRvD...1..467C,1970PhRvD...2.2103C}, using the same method Cheng and Wu further calculated other cross-sections for high-energy photon-photon scattering to double pion, muon and electron--positron pairs:  \\begin{itemize}  \\item the process of double muon pair production $\\gamma_{1}+\\gamma_{2}\\rightarrow \\mu^{+}+\\mu^{-}+ \\mu^{+}+\\mu^{-}$ and its cross-section can be obtained by replacing $m_e\\rightarrow m_\\mu$ in Eq.~(\\ref{cwresult}), thus $\\sigma_{\\gamma\\gamma}(s)\\sim 1.5\\cdot 10^{-4} \\mu b$.  \\item the process of double pion pair production $\\gamma_{1}+\\gamma_{2}\\rightarrow \\pi^{+}+\\pi^{-}+ \\pi^{+}+\\pi^{-}$ and its extremely small cross-section \\begin{equation}\\label{cwpi} \\lim_{s\\rightarrow \\infty}\\sigma_{\\gamma\\gamma}(s)= \\frac{\\alpha^4}{144\\pi m_\\pi^2}[7\\zeta(3)+10]\\sim 0.23\\cdot 10^{-5} \\mu b, \\end{equation}  \\item the process of one pion pair and one electron--positron pair production $\\gamma_{1}+\\gamma_{2}\\rightarrow e^{+}+e^{-}+ \\pi^{+}+\\pi^{-}$ and its small cross-section \\begin{equation}\\label{cwpiee} \\lim_{s\\rightarrow \\infty}\\sigma_{\\gamma\\gamma}(s)= \\frac{2\\alpha^4}{27\\pi m_\\pi^2}\\left[\\left(\\ln\\frac{m_\\pi^2}{m_e^2}\\right)^2 +\\frac{16}{3}\\ln\\frac{m_\\pi^2}{m_e^2}+\\frac{163}{18}\\right]\\sim 0.26\\cdot 10^{-3} \\mu b, \\end{equation} which is more than one hundred times larger than (\\ref{cwpi}).  \\item the process of one pion pair and one muon-antimuon pair production $\\gamma_{1}+\\gamma_{2}\\rightarrow \\mu^{+}+\\mu^{-}+ \\pi^{+}+\\pi^{-}$ and its small cross-section can be obtained by replacing $m_e\\rightarrow m_\\mu$ in Eq.~(\\ref{cwpiee}) everywhere.  \\end{itemize}  \\subsection{Electron-nucleus bremsstrahlung and pair production by a photon in the field of a nucleus}\\label{bremss}  The other two important QED processes, related by the crossing symmetry are the electron-nucleus bremsstrahlung (\\ref{eibr}) and creation of electron--positron pair by a photon in the field of a nucleus (\\ref{gi2p}). These processes were considered already in the early years of QED. They are of higher order, comparing to the Compton scattering and the Breit--Wheeler processes, respectively, and contain one more vertex connecting fermions with the photon.  The nonrelativistic cross-section for the process (\\ref{eibr}) was derived by Sommerfeld \\cite{1931AnP...403..257S}. Here we remind the basic results obtained in the relativistic case by Bethe and Heitler \\cite{1934RSPSA.146...83B} and independently by Sauter \\cite{1934AnP...412..404S}. The Feynman diagram for bremsstrahlung can be imagined considering for the Compton scattering, but treating one of the photons as virtual one corresponding to an external field. We consider this process in Born approximation, and the momentum recoil of the nucleus is neglected. Integrating the differential cross-section over all directions of the photon and the outgoing electron one has, see e.g. \\cite{1982els..book.....B} \\begin{equation} \\begin{array}{lr} d\\sigma=Z^2\\alpha r_e^2\\frac{d\\omega}{\\omega}\\frac{p^\\prime}{p}\\left\\{\\frac{4}{3}-2\\epsilon\\epsilon^\\prime\\frac{p^2+p^{\\prime2}}{p^2 p^{\\prime2}}+m_e^2\\left(l\\frac{\\epsilon^\\prime}{p^3}+l^\\prime\\frac{\\epsilon}{p^{\\prime3}-\\frac{ll^\\prime}{pp^\\prime}}\\right)+\\right.  \\\\ \\left.+L\\left[\\frac{8\\epsilon\\epsilon^\\prime}{3pp^\\prime}+\\frac{\\omega^2}{p^3p^{\\prime3}}\\left(\\epsilon^2\\epsilon^{\\prime2}+p^2p^{\\prime2}+m_e^2\\epsilon\\epsilon^\\prime\\right)+\\frac{m_e^2\\omega}{2pp^\\prime}\\left(l\\frac{\\epsilon\\epsilon^\\prime+p^2}{p^3}-l^\\prime\\frac{\\epsilon\\epsilon^\\prime+p^{\\prime2}}{p^{\\prime3}}\\right)\\right]\\right\\}, \\end{array} \\end{equation} where \\begin{equation} L=\\log\\frac{\\epsilon\\epsilon^\\prime+pp^\\prime-m_e^2}{\\epsilon\\epsilon^\\prime-pp^\\prime-m_e^2}, \\quad l=\\log\\frac{\\epsilon+p}{\\epsilon-p}, \\quad l^\\prime=\\log\\frac{\\epsilon^\\prime+p^\\prime}{\\epsilon^\\prime-p^\\prime}, \\end{equation} and $\\omega$ is photon energy, $p$ and $p^\\prime$ are electron momenta before and after the collision, respectively, $\\epsilon$ and $\\epsilon^\\prime$ are its initial and final energies.  The averaged cross-section for the process of pair production by a photon in the field of a nucleus may be obtained by applying the transformation rules relating the processes (\\ref{gi2p}) and (\\ref{eibr}), see e.g. \\cite{1982els..book.....B}. The result is \\begin{equation} \\begin{array}{lcr} d\\sigma = Z^2\\alpha r_e^2\\frac{p_+p_-}{\\omega^3}d\\epsilon_+\\biggl\\{-\\frac{4}{3}-2\\epsilon_+\\epsilon_-\\frac{p_+^2+p_-^2}{p_+^2p_-^2}+m_e^2\\left(l_-\\frac{\\epsilon_+}{p_-^3}+l_+\\frac{\\epsilon_-}{p_+^3-\\frac{l_+l_-}{p_+p_-}}\\right)+  \\\\ +L\\biggl[\\frac{8\\epsilon_+\\epsilon_-}{3p_+p_-}+\\frac{\\omega^2}{p_+^3p_-^3}\\left(\\epsilon_+^2\\epsilon_-^2+p_+^2p_-^2-m_e^2\\epsilon_+\\epsilon_-\\right)- \\\\ -\\frac{m_e^2\\omega}{2p_+p_-}\\left(l_+\\frac{\\epsilon_+\\epsilon_- -p_+^2}{p_+^3}+l_-\\frac{\\epsilon_+\\epsilon_- -p_-^2}{p_-^3}\\right)\\biggr]\\biggr\\}, \\label{dsigmapm} \\end{array} \\end{equation} where \\begin{equation} L=\\log\\frac{\\epsilon_+\\epsilon_-+p_+p_-+m_e^2}{\\epsilon_+\\epsilon_- -p_+p_-+m_e^2}, \\quad l_\\pm=\\log\\frac{\\epsilon_\\pm+p_\\pm}{\\epsilon_\\pm-p_\\pm}, \\end{equation} and $p_\\pm$ and $\\epsilon_\\pm$ are momenta and energies of positron and electron respectively.  The total cross-section for this process is given in Section \\ref{higher}, see (\\ref{sigmaZ1Z2}). In the ultrarelativistic approximation relaxing the condition $Z\\alpha\\ll1$ the pair production by the process (\\ref{gi2p}) was treated by Bethe and Maximon in \\cite{1954PhRv...93..768B,1954PhRv...93..788D}. The total cross-section (\\ref{sigmaZ1Z2}) is becoming then \\begin{equation} \\sigma=\\frac{28}{9}Z^2\\alpha r_e^2\\left(\\log\\frac{2\\omega}{m_e}-\\frac{109}{42}-f(Z\\alpha)\\right), \\label{sigmaZ1Z2r} \\end{equation} where \\begin{equation} f(Z\\alpha)=\\gamma_E+Re\\Psi(1+iZ\\alpha)=(Z\\alpha)^2 \\sum_{n=1}^\\infty\\frac{1}{n[n^2+(Z\\alpha)^2]}. \\end{equation}  \\subsection{Pair production in collision of two ions}\\label{ionion}  The process of the pair production by two ions (\\ref{2pe+e-}) in ultrarelativistic approximation was considered by Landau and Lifshitz \\cite{Landau1934} and Racah \\cite{Racah1937}. For the modern review of these topics see \\cite{2007PhR...453....1B}. The corresponding differential cross-section with logarithmic accuracy can be obtained from the differential cross-section (\\ref{dsigmapm}) taking its ultrarelativistic approximation $\\gamma\\gg1$ for the Lorentz factor of the relative motion of the two nuclei with charges $Z_1$ and $Z_2$ respectively, and treating the real photon line in the process (\\ref{gi2p}) as a virtual photon corresponding to the external field of the nucleus. One should then multiply the cross-section by the spectrum of these equivalent photons, see Section \\ref{higher}, and the result is \\begin{equation} d\\sigma=\\frac{8}{\\pi}r_e^2(Z_1Z_2\\alpha)^2\\frac{d\\epsilon_+d\\epsilon_-}{(\\epsilon_++\\epsilon_-)^4}\\left(\\epsilon_+^2+\\epsilon_-^2+\\frac{2}{3}\\epsilon_+\\epsilon_-\\right)\\log\\frac{\\epsilon_+\\epsilon_-}{m_e(\\epsilon_++\\epsilon_-)}\\log\\frac{m_e\\gamma}{(\\epsilon_++\\epsilon_-)}. \\end{equation}  The total cross-section is given in Section \\ref{higher}, see (\\ref{sigmaL}) and (\\ref{sigmaR}). More recent results containing higher order in $\\alpha$ corrections are obtained in \\cite{2002PhLB..538...45B,2002PhRvA..66d2720B}, see also \\cite{2006PPNL....3..246V,2007PPNL....4...18V}.  \\begin{figure}[!ht] \\centering \\includegraphics[width=5cm]{coulcor.eps} \\caption{Classification of the $e^+e^-$ pair production by the number of photons attached to a nucleus. Reproduced from \\cite{2007PhR...453....1B}.} \\label{coulcor} \\end{figure} Lepton pair production in relativistic ion collision to all orders in $Z\\alpha$ with logarithmic accuracy is studied in \\cite{2003JPhG...29.1227G} where the matrix elements are separated in different classes, see Fig. \\ref{coulcor}, according to numbers of photon lines attached to a given nucleus \\begin{equation} M=M_{(\\mathrm{i})}+M_{(\\mathrm{ii})}+M_{(\\mathrm{iii})}+M_{(\\mathrm{iv})}. \\end{equation} The Born amplitude $\\left|M_{(\\mathrm{i})}\\right|^2$ corresponding to the lowest order in $Z\\alpha$ with one photon line attached to each nucleus was computed by Landau and Lifshitz \\cite{Landau1934} who obtained the famous $L_\\gamma^3$ dependence of the cross-section in (\\ref{sigmaL}). It should be mentioned that the Racah formula (\\ref{sigmaR}) in contrast with the Landau and Lifshitz result (\\ref{sigmaL}) contains also terms proportional to $L_\\gamma^2$ and $L_\\gamma$. These terms come from the absolute square of $M_{(\\mathrm{ii})}$ and $M_{(\\mathrm{iii})}$ and their interference with the Born amplitude $M_{(\\mathrm{i})}$ \\cite{2002PhRvA..66d2720B}. Their result for the Coulomb corrections which is defined as the difference between the full cross-section and the Born approximation, in order $L_\\gamma^2$ is of the ``Bethe--Maximon'' type \\cite{2002PhRvA..66d2720B} \\begin{equation} \\sigma_\\mathrm{C}=\\frac{28}{27\\pi}r_e^2(Z_1Z_2\\alpha)^2\\left[f(Z_1\\alpha)f(Z_2\\alpha)\\right](L_\\gamma^2+{\\mathcal O}(L_\\gamma)). \\label{sigmaC} \\end{equation} The Coulomb corrections here are up to $L_\\gamma^2$-term in the Racah formula (\\ref{sigmaR}).  The calculations mentioned above were all made as early as in the 1930s. Clearly, at that time only $e^+e^-$ pair production was discussed. However, these calculations can be considered for any lepton pair production, for example $\\mu^+\\mu^-$ pair production, as long as the total energy in the center of mass of the system is large enough. However, simple substitution $m_e\\longrightarrow m_l$, where $l$ stands for any lepton is not sufficient since when the energy reaches the inverse radius $\\Lambda=1/R\\sim30{\\mathrm MeV}$ of the nucleus the electric field of the nucleus cannot be approximated as a Coulomb field of a point-like particle \\cite{1998PhRvD..57.4025I}. The review of computation for lepton pair production can be found in \\cite{1973SvPhU..16..322G} and \\cite{1975PhR....15..181B}.  Another effect of large enough collision energy is multiple pair production. Early work on this subject started with the observation that the impact-parameter-dependent total pair production probability computed in the lowest order perturbation theory is larger than one. The analysis of Ref. \\cite{2002PhLB..538...45B} devoted to the study of the corresponding Feynman diagrams in the high-energy limit leads to the probability of $N$ lepton pair production obeying the Poisson distribution. For a review of this topic see \\cite{2007PhR...453....1B}.  Two photon particle pair production by collision of two electrons or electron and positron (see Fig. \\ref{twophoton}, where 4-momenta $p_1$ and $p_2$ correspond to their momenta) were studied in storage rings in Novosibirsk ($e^+e^-\\rightarrow e^+e^-e^+e^-$ \\cite{1971PhLB...34..663B}) and in Frascati at ADONE ($e^+e^-\\rightarrow e^+e^-e^+e^-$ \\cite{Bacci1972}, $e^+e^-\\rightarrow e^+e^-\\mu^+\\mu^-$ \\cite{1974PhRvL..32..385B}, $e^+e^-\\rightarrow e^+e^-\\pi^+\\pi^-$ \\cite{1974PhLB...48..380O}), see also \\cite{1980LNP...134....9B}. At high energy the total cross-section of two photon production of lepton pairs is given by \\cite{Landau1934,1975PhR....15..181B} \\begin{equation} \\sigma_{e^\\pm e^-\\rightarrow e^\\pm e^- l^+l^-}=\\frac{28\\alpha^4}{27\\pi m_l^2}\\left(\\ln\\frac{s}{m_e^2}\\right)^2\\ln\\frac{s}{m_e^2};\\quad l\\equiv e \\;\\; {\\mathrm or} \\;\\; \\mu. \\end{equation} see c.f. Eq. (\\ref{sigmaL}).  \\subsection{QED description of pair production}\\label{qedpair}  We turn now to a Sauter-Heisenberg-Euler process in QED. An external electromagnetic field is incorporated by adding to the quantum field $A_\\mu$ in (\\ref{L0int}) an unquantized external vector potential $A^{\\rm e}_\\mu$, so that the total interaction becomes \\begin{equation} {\\mathcal L}_{\\rm int}+ {\\mathcal L}^{\\rm e}_{\\rm int} = -e\\bar\\psi(x)\\gamma^\\mu \\psi(x) \\left[ A_\\mu(x)+ A^{\\rm e}_\\mu(x)\\right] . \\label{intc} \\end{equation} Instead of an operator formulation, one can derive the quantum field theory from a functional integral formulation, see e.g. \\cite{Kleinert2004}, in which the quantum mechanical partition function is described by \\begin{equation} Z[A^{\\rm e}]=\\int [{\\mathcal D}\\psi {\\mathcal D}\\bar\\psi {\\mathcal D}A_\\mu] \\exp \\left[ i\\int d^4x ({\\mathcal L} + {\\mathcal L}^{\\rm e}_{\\rm int} ) \\right], \\label{path} \\end{equation} to be integrated over all fluctuating electromagnetic and Grassmannian electron fields. The normalized quantity $Z[A^{\\rm e}]$ gives the amplitude for the vacuum to vacuum transition in the presence of the external classical electromagnetic field: \\begin{equation} \\langle {\\rm out}, 0|0, {\\rm in}\\rangle = \\frac{Z[A^{\\rm e}]}{ Z[0]}, \\label{vvamplitude} \\end{equation} where $|0, {\\rm in}\\rangle$ is the initial vacuum state at the time $t=t_-\\rightarrow -\\infty$, and $\\langle {\\rm out}, 0|$ is the final vacuum state at the time $t=t_+\\rightarrow +\\infty$. By selecting only the one-particle irreducible Feynman diagrams in the perturbation expansion of $Z[A^{\\rm e}]$ one obtains the effective action as a functional of $A^{\\rm e}$: \\begin{equation} \\Delta{\\mathcal A}_{\\rm eff}[A^{\\rm e}]\\equiv -i\\ln \\langle {\\rm out}, 0|0, {\\rm in}\\rangle. \\label{eaction} \\end{equation} In general, there exists no local effective Lagrangian density $\\Delta{\\mathcal L}_{\\rm eff}$ whose space-time integral is $\\Delta{\\mathcal A}_{\\rm eff}[A^{\\rm e}]$. An infinite set of derivatives would be needed, i.e., $\\Delta{\\mathcal L}_{\\rm eff}$ would have the arguments $A^{\\rm e}(x),\\partial_\\mu A^{\\rm e}(x), \\partial_\\mu\\partial _\\nu A^{\\rm e}(x), \\dots~$, containing gradients of arbitrary high order. With presently available methods it is possible to calculate a few terms in such a gradient expansion, or a semi-classical approximation {\\it \\`a l\\`a} JWKB for an arbitrary but smooth space-time dependence (see Section~3.21ff in Ref.~\\cite{Kleinert2004}). Under the assumption that the external field $A^{\\rm e}(x)$ varies smoothly over a finite space-time region, we may define an approximately local effective Lagrangian $\\Delta{\\mathcal L}_{\\rm eff}[A^{\\rm e}(x)]$, \\begin{equation} \\Delta{\\mathcal A}_{\\rm eff}[A^{\\rm e}]\\simeq \\int d^4x \\Delta{\\mathcal L}_{\\rm eff}[A^{\\rm e}(x)] \\approx V\\Delta t \\Delta{\\mathcal L}_{\\rm eff}[A^{\\rm e}], \\label{effl} \\end{equation} where $V$ is the spatial volume and time interval $\\Delta t=t_+-t_-$.  For a large time interval $\\Delta t=t_+-t_-\\rightarrow \\infty$, the amplitude of the vacuum to vacuum transition (\\ref{vvamplitude}) has the form, \\begin{equation} \\langle {\\rm out}, 0|0 ,{\\rm in}\\rangle = e^{-i(\\Delta{\\mathcal E}_0-i\\Gamma/2)\\Delta t}, \\label{vvamplitude1} \\end{equation} where $\\Delta{\\mathcal E}_0={\\mathcal E}_0(A^{\\rm e})-{\\mathcal E}_0(0)$ is the difference between the vacuum energies in the presence and the absence of the external field, $\\Gamma$ is the vacuum decay rate, and $\\Delta t$ the time over which the field is nonzero. The probability that the vacuum remains as it is in the presence of the external classical electromagnetic field is \\begin{equation} |\\langle {\\rm out}, 0|0 ,{\\rm in}\\rangle|^2 = e^{-2{\\rm Im}\\Delta{\\mathcal A}_{\\rm eff}[A^{\\rm e}]}. \\label{probability} \\end{equation} This determines the decay rate of the vacuum in an external electromagnetic field: \\begin{equation} \\frac{ \\Gamma}{V}= \\frac{2{\\rm \\,Im}\\Delta{\\mathcal A}_{\\rm eff}[A^{\\rm e}]}{V\\Delta t} \\approx 2{\\rm \\,Im}\\Delta{\\mathcal L}_{\\rm eff}[A^{\\rm e}]. \\label{path21} \\end{equation} The vacuum decay is caused by the production of electron and positron pairs. The external field changes the energy density by \\begin{equation} \\frac{ \\Delta{\\mathcal E}_0}{V}= -\\frac{{\\rm \\,Re}\\Delta{\\mathcal A}_{\\rm eff}[A^{\\rm e}]}{V\\Delta t} \\approx -{\\rm \\,Re}\\Delta{\\mathcal L}_{\\rm eff}[A^{\\rm e}]. \\label{path21e} \\end{equation}  \\subsubsection{Schwinger formula for pair production in uniform electric field}\\label{Schwingerformula}  The Dirac fields appears quadratically in the partition functional (\\ref{path}) and can be integrated out, leading to \\begin{equation} Z[A^{\\rm e}]=\\int  {\\mathcal D}A_\\mu\\, {\\rm Det} \\{i\\!\\!\\not{\\!\\partial}-e[\\not{\\hspace{-5pt}A(x)}+ \\not{\\hspace{-5pt}A}^{\\rm e}(x)] -m_e+i\\eta\\};\\quad \\not{\\!\\partial}\\equiv\\gamma^\\mu\\partial_\\mu, \\quad \\not{\\!A}\\equiv\\gamma^\\mu A_\\mu, \\end{equation} where Det denotes the functional determinant of the Dirac operator. Ignoring the fluctuations of the electromagnetic field, the result is a functional of the external vector potential  $A^{\\rm e}(x)$: \\begin{equation} Z[A^{\\rm e}]\\approx {\\rm const}\\times {\\rm Det} \\{i\\!\\!\\not{\\!\\partial}-e \\not{\\hspace{-5pt}A}^{\\rm e}(x)-m_e+i\\eta\\}. \\label{path1} \\end{equation} The infinitesimal constant $i\\eta$ with $ \\eta >0$ specifies the treatment of singularities in energy integrals. From Eqs.~(\\ref{vvamplitude})--(\\ref{path1}), the effective action (\\ref{probability}) is given by \\begin{equation} \\Delta{\\mathcal A}_{\\rm eff}[A^{\\rm e}]=-i {\\rm Tr}\\ln \\left\\{\\left[ i\\!\\!\\not{\\!\\partial}- e\\not{\\hspace{-5pt}A}^{\\rm e}(x)-m_e+i\\eta\\right] \\frac{1}{ i\\!\\!\\not{\\!\\partial} -m_e+i\\eta}\\right\\}, \\label{eaction12} \\end{equation} where Tr denotes the functional and Dirac trace. In physical units, this is of order $\\hbar$. The result may be expressed as a one-loop Feynman diagram, so that one speaks of one-loop approximation. More convenient will be the equivalent expression \\begin{equation} \\Delta{\\mathcal A}_{\\rm eff}[A^{\\rm e}]= -\\frac{i}{2}\\,{\\rm Tr}\\ln \\left(\\{[i\\!\\!\\not{\\!\\partial}- e\\not{\\hspace{-5pt}A}^{\\rm e}(x)]^2-m_e^2+i\\eta\\} \\frac{1}{ -\\partial ^2-m_e^2+i\\eta}\\right), \\label{eaction1} \\end{equation} where \\begin{equation} [i\\!\\!\\not{\\!\\partial}- e\\not{\\hspace{-5pt}A}^{\\rm e}(x)]^2= [i\\partial_\\mu-eA^{\\rm e}_\\mu(x)]^2+\\frac{e}{2}\\sigma^{\\mu\\nu}F^{\\rm e}_{\\mu\\nu}, \\label{sqrt} \\end{equation} where $\\sigma^{\\mu\\nu}\\equiv \\frac{i}{2}[\\gamma^\\mu,\\gamma^\\nu]$,\\, $F^{\\rm e}_{\\mu\\nu}=\\partial_\\mu A^{\\rm e}_\\nu-\\partial_\\nu A^{\\rm e}_\\mu$. Using the identity \\begin{equation} \\ln\\frac{a_2}{ a_1}=\\int_0^\\infty \\frac{ds}{ s}\\big[ e^{is(a_1+i\\eta)}-e^{is(a_2+i\\eta)}\\big], \\label{identityab} \\end{equation} Eq.~(\\ref{eaction1}) becomes \\begin{equation} \\Delta{\\mathcal A}_{\\rm eff}[A^{\\rm e}] =\\frac{i}{2}\\int_0^\\infty \\frac{ds}{ s}e^{-is(m_e^2-i\\eta)}{\\rm Tr} \\langle x| e^{ is\\left\\{[i\\partial_\\mu-eA^{\\rm e}_\\mu(x)]^2+\\frac{e}{2}\\sigma^{\\mu\\nu}F^{\\rm e}_{\\mu\\nu}\\right\\}} -e^{-is\\partial ^2}|x\\rangle , \\label{probability01} \\end{equation} where $\\langle x|\\{\\cdot\\cdot\\cdot\\}|x\\rangle $ are the diagonal matrix elements in the local basis $|x\\rangle $. This is defined by the matrix elements with the momentum eigenstates $|k\\rangle$ being plane waves: $\\langle x|k\\rangle= e^{-ikx}$. The symbol ${\\rm Tr}$ denotes integral $\\int d^4x$ in space-time and the trace in spinor space. For constant electromagnetic fields, the integrand in (\\ref{probability01}) does not depend on $x$, and $\\sigma^{\\mu\\nu}F^{\\rm e}_{\\mu\\nu}$ commutes with all other operators. This will allows us to calculate the exponential in Eq.~(\\ref{probability01}) explicitly. The presence of $-i \\eta $ in the mass term ensures convergence of integral for $s\\rightarrow\\infty$.   If only a constant electric field ${\\bf E}$ is present, it may be assumed to point along the ${\\hat{\\bf z}}$-axis, and one can choose a gauge such that $A^{\\rm e}_z=-Et$ is the only nonzero component of $A^{\\rm e}_\\mu$. Then one finds \\begin{equation} {\\rm tr}\\exp is\\left[\\frac{e}{2}\\sigma^{\\mu\\nu}F^{\\rm e}_{\\mu\\nu}\\right]=4\\cosh(seE), \\label{iz1} \\end{equation} where the symbol ${\\rm tr}$ denotes the trace in spinor space. Using the commutation relation $[\\partial _0,x^0]=1$, where $x^0=t$, one computes the exponential term in the effective action (\\ref{probability01}) (c.e.g.~\\cite{Itzykson2006}) \\begin{equation} \\langle x| \\exp is\\left[(i\\partial _\\mu-eA^{\\rm e}_\\mu(x))^2+\\frac{e}{2} \\sigma^{\\mu\\nu}F^{\\rm e}_{\\mu\\nu}\\right] |x\\rangle =\\frac{eE}{ (2\\pi)^2is}\\coth(eEs). \\label{iz2} \\end{equation} The second term in Eq.~(\\ref{probability01}) is obtained by setting $E= 0$ in Eq.~(\\ref{iz2}), so that the effective action (\\ref{probability01}) yields, \\begin{equation} \\Delta{\\mathcal A}_{\\rm eff}= \\frac{1}{ 2(2\\pi)^2}\\int d^4x \\int_0^\\infty \\frac{ds}{ s^3} \\left[eEs\\coth(eEs)-1\\right] e^{-is(m^2_e-i\\eta)}. \\label{oneloops1} \\end{equation} Since the field is constant, the integral over $x$ gives a volume factor, and the effective action (\\ref{probability01}) can be attributed to the space-time integral over an effective Lagrangian (\\ref{effl}) \\begin{equation} \\Delta{\\mathcal L}_{\\rm eff} = \\frac{1}{ 2(2\\pi)^2}\\int_0^\\infty \\frac{ds}{ s^3} \\left[eEs\\coth(eEs)-1\\right] e^{-is(m^2_e-i\\eta)}. \\label{oneloopl1us} \\end{equation}  By expanding the integrand in Eq.~(\\ref{oneloopl1us}) in powers of $e$, one obtains, \\begin{equation} \\frac{1}{ s^3}\\left[eEs\\coth(eEs)-1\\right] e^{-i s(m^2_e-i\\eta)}= \\left[\\frac{e^2}{ 3s}E^2-\\frac{e^4s}{ 45}E^4 +{\\mathcal O}(e^6)\\right] e^{-i s(m^2_e-i\\eta)}. \\label{oneloopl1s20} \\end{equation} The small-$s$ divergence in the integrand, \\begin{equation} \\frac{e^2}{ 3}E^2 \\frac{1}{ 2(2\\pi)^2}\\int_0^\\infty \\frac{ds}{ s}e^{-is(m^2_e-i\\eta)}, \\label{oneloopl1s2} \\end{equation} is proportional to the electric term in the original Maxwell Lagrangian. The divergent term (\\ref{oneloopl1s2}) can therefore be removed by a renormalization of the field $E$. Thus we subtract a counterterm in Eq.~(\\ref{oneloopl1us}) and form \\begin{equation} \\Delta{\\mathcal L}_{\\rm eff} = \\frac{1}{ 2(2\\pi)^2}\\int_0^\\infty \\frac{ds}{ s^3} \\left[eEs\\coth(eEs)-1-\\frac{e^2}{ 3}E^2s^2\\right] e^{-is(m^2_e-i\\eta)}. \\label{oneloopl1} \\end{equation} Remembering Eq.~(\\ref{path21}), we find from (\\ref{oneloopl1})  the decay rate of the vacuum per unit volume \\begin{equation} \\frac{ \\Gamma }{V} =\\frac{1}{ (2\\pi)^2}{\\rm Im}\\int_0^\\infty \\frac{ds}{ s^3}\\left[eEs\\coth(eEs)-1-\\frac{e^2}{ 3}E^2s^2\\right] e^{-is(m_e^2-i\\eta)}. \\label{gprobability} \\end{equation}  The integral (\\ref{gprobability}) can be evaluated analytically by the method of residues. Since the integrand is even, the integral can be extended to the entire $s$-axis. After this, the integration contour is deformed to enclose the negative imaginary axis and to pick up the contributions of the poles of the $\\coth$ function at $s=n\\pi/eE$. The result is \\begin{equation} \\!\\!\\!\\!\\!\\!\\!\\! \\frac{ \\Gamma }{V} =\\frac{\\alpha E^2}{ \\pi^2}\\sum_{n=1}^\\infty \\frac{1}{ n^2}\\exp \\left(-\\frac{n\\pi E_c}{ E}\\right). \\label{probability1} \\end{equation} This result, i.e. the {\\it Schwinger formula} \\cite{1951PhRv...82..664S,1954PhRv...93..615S,1954PhRv...94.1362S} is valid to lowest order in $\\hbar$ for arbitrary constant electric field strength.  An analogous calculation for a charged scalar field yields \\begin{equation} \\frac{ \\Gamma }{V} =\\frac{\\alpha E^2}{ 2\\pi^2}\\sum_{n=1}^\\infty \\frac{(-1)^{n+1}}{ n^2}\\exp \\left(-\\frac{n\\pi E_c}{ E}\\right), \\label{bosonrate} \\end{equation} which generalizes the Weisskopf treatment being restricted to the leading term $n=1$. These Schwinger results complete the derivation of the probability for pair productions. The leading $n=1$ -terms of (\\ref{probability1}) and (\\ref{bosonrate}) agrees with the JWKB results we discuss in Section \\ref{semi}, and thus the correct Sauter exponential factor (\\ref{transmission}) and Heisenberg-Euler imaginary part of the effective Lagrangian (\\ref{herate}).  Narozhny and Nikishov \\cite{Narozhnyi:1970uv} have expressed Eq.~(\\ref{probability}) through the probability of one pair production $P_1$, of $n$ pair production $P_n$ with $n=1,2,3,\\cdot\\cdot\\cdot$ as well as the average number of pair productions \\begin{equation} |\\langle {\\rm out}, 0|0 ,{\\rm in}\\rangle|^2 = 1- P_1-P_2-P_3-\\cdot\\cdot\\cdot , \\label{nprobability} \\end{equation} where $P_n, (n=1,2,3,\\cdot\\cdot\\cdot)$ is the probability of $n$ pair production, and the probability of one pair production is, \\begin{equation} P_1 = V\\Delta t\\frac{\\alpha E^2}{2\\pi^2}\\ln \\left(1-e^{-\\frac{\\pi m_e^2}{eE}}\\right) e^{-2V\\Delta t {\\rm Im}\\Delta{\\mathcal L}_{\\rm eff}[A^{\\rm e}]} . \\label{nprobability1} \\end{equation} The average number $\\bar N$ of pair productions is then given by \\begin{equation} \\bar N=\\sum_{n=1}^\\infty nP_n =V\\Delta t \\frac{\\alpha E^2}{ \\pi^2}\\exp \\left(-\\frac{\\pi E_c}{ E}\\right), \\label{nprobabilitym1} \\end{equation} which is the quantity directly related to the experimental measurements.  \\subsubsection{Pair production in constant electromagnetic fields}\\label{EandB}  Since the QED theory is gauge and Lorentz invariant, effective action $\\Delta{\\mathcal A}_{\\rm eff}$ and Lagrangian $\\Delta{\\mathcal L}_{\\rm eff}$ are expressed as functionals of the scalar and pseudoscalar invariants $S,P$ (\\ref{lightlike}). Thus they must be invariant under the discrete duality transformation: \\begin{equation} |{\\bf B}|\\rightarrow -i|{\\bf E}|,\\quad |{\\bf E}|\\rightarrow i|{\\bf B}|, \\label{duality1} \\end{equation} i.e., \\begin{equation} \\beta\\rightarrow -i\\varepsilon ,\\quad \\varepsilon\\rightarrow i\\beta. \\label{duality2} \\end{equation} This implies that in the case ${\\bf E}=0$ and ${\\bf B}\\not= 0$, results can be simply obtained by replacing $|{\\bf E}|\\rightarrow i|{\\bf B}|$ in Eqs.~(\\ref{iz2}), (\\ref{oneloopl1}), (\\ref{gprobability}): \\begin{equation} \\langle x| \\exp is\\left[(i\\partial _\\mu-eA^{\\rm e}_\\mu(x))^2+\\frac{e}{2}\\sigma^{\\mu\\nu}F^{\\rm e}_{\\mu\\nu}\\right] |x\\rangle =\\frac{eB}{ (2\\pi)^2is}\\cot(eBs), \\label{iz2b} \\end{equation} and \\begin{equation} \\Delta{\\mathcal L}_{\\rm eff} = \\frac{1}{ 2(2\\pi)^2}\\int_0^\\infty \\frac{ds}{ s^3} \\left[eBs\\cot(eBs)-1+\\frac{e^2}{ 3}B^2s^2\\right] e^{-is(m^2_e-i\\eta)}. \\label{oneloopl1b} \\end{equation}  In the presence of both constant electric and magnetic fields ${\\bf E}$ and ${\\bf B}$, we adopt parallel ${\\bf E}_{\\rm CF}$ and ${\\bf B}_{\\rm CF}$ pointing along the ${\\hat{\\bf z}}$-axis in the {\\it center-of-fields frame}, as discussed after Eqs.~(\\ref{eblandau}), (\\ref{elorentz}), (\\ref{blorentz}). We can choose a gauge such that only $A^{\\rm e}_z=-E_{\\rm CF}t$, $A^{\\rm e}_y=B_{\\rm CF}x^1$ are nonzero. Due to constant fields, the exponential in the effective action Eq.~(\\ref{probability01}) can be factorized into a product of the magnetic part and the electric part. Following the same method used to compute the electric part (\\ref{iz1},\\ref{iz2}), one can compute the magnetic part by using the commutation relation $[\\partial _1,x^1]=1$, where $x^1=x$. Or one can make the substitution (\\ref{duality1}) to obtain the magnetic part, based on the discrete symmetry of duality. As results, Eqs.~(\\ref{iz1}), (\\ref{iz2}) become \\begin{equation} {\\rm tr}\\exp is\\left[\\frac{e}{2}\\sigma^{\\mu\\nu}F^{\\rm e}_{\\mu\\nu}\\right] =4\\cosh(seE_{\\rm CF})\\cos(seB_{\\rm CF}), \\label{iz1eb} \\end{equation} and \\begin{align} &\\langle x| \\exp is\\left\\{[i\\partial _\\mu-eA^{\\rm e}_\\mu(x)]^2+\\frac{e}{2}\\sigma^{\\mu\\nu}F^{\\rm e}_{\\mu\\nu}\\right\\} |x\\rangle \\nonumber\\\\ &=\\frac{1}{(2\\pi)^2}\\frac{eE_{\\rm CF}}{ is}\\coth(seE_{\\rm CF}) \\frac{eB_{\\rm CF}}{ s}\\cot(seB_{\\rm CF}). \\label{iz2eb} \\end{align} In this special frame, the effective Lagrangian is then given by \\begin{align} \\Delta{\\mathcal L}_{\\rm eff}&= \\frac{1}{ 2(2\\pi)^2}\\int_0^\\infty \\frac{ds}{ s^3} \\Big[e^2E_{\\rm CF}B_{\\rm CF} s^2\\coth(seE_{\\rm CF} ) \\cot(seB_{\\rm CF})\\nonumber\\\\ &-1-\\frac{e^2}{ 3}(E^2_{\\rm CF}-B^2_{\\rm CF})s^2 \\Big] \\cdot e^{-is(m^2_e-i\\eta)}. \\label{onelooplcf} \\end{align} Using definitions in Eqs.~(\\ref{lightlike}), (\\ref{ab}), (\\ref{fieldinvariant}), we obtain the effective Lagrangian \\begin{align} \\Delta{\\mathcal L}_{\\rm eff} &\\!=\\! \\frac{1}{ 2(2\\pi)^2}\\int_0^\\infty \\frac{ds}{ s^3} \\Big[e^2\\varepsilon\\beta s^2\\coth(e\\varepsilon s )\\cot(e\\beta s )\\nonumber\\\\ &-1-\\frac{e^2}{ 3}(\\varepsilon^2-\\beta^2)s^2\\Big] e^{-is(m^2_e-i\\eta)}; \\label{oneloopl} \\end{align} and the decay rate \\begin{align} \\frac{ \\Gamma }{V}&=\\frac{1}{ (2\\pi)^2}{\\rm Im}\\int_0^\\infty \\frac{ds}{ s^3} \\Big[e^2\\varepsilon\\beta s^2\\coth(e\\varepsilon s )\\cot(e\\beta s )\\nonumber\\\\ &-1-\\frac{e^2}{ 3}(\\varepsilon^2-\\beta^2)s^2\\Big] e^{-is(m_e^2-i\\eta)}, \\label{gprobabilityeb} \\end{align} in terms of the invariants $\\varepsilon$ and $\\beta$ (\\ref {fieldinvariant}) for arbitrary electromagnetic fields ${\\bf E}$ and ${\\bf B}$.  The integral (\\ref{gprobabilityeb}) is evaluated as in Eq.~(\\ref{probability1}) by the method of residues, and yields \\cite{1951PhRv...82..664S,1954PhRv...93..615S,1954PhRv...94.1362S} \\begin{equation} \\frac{ \\Gamma }{V}=\\frac{  \\alpha   \\varepsilon^2}{ \\pi^2 } \\sum_{n=1}  \\frac{1}{n^2} \\frac{ n\\pi\\beta / \\varepsilon } {\\tanh {n\\pi \\beta/ \\varepsilon}}\\exp\\left(-\\frac{n\\pi E_c}{ \\varepsilon}\\right), \\label{probabilityeh} \\end{equation} which reduces for $\\beta \\rightarrow 0$ (${\\bf B}=0$) correctly to (\\ref{probability1}). The $n=1$ -term is the JWKB approximation (\\ref{wkbehfermion2}).  The analogous result for bosonic fields is \\begin{equation} \\frac{ \\Gamma }{V}=\\frac{  \\alpha   \\varepsilon^2}{ 2\\pi^2 } \\sum_{n=1}  \\frac{(-1)^n}{n^2} \\frac{ n\\pi\\beta / \\varepsilon } {\\sinh {n\\pi \\beta/ \\varepsilon}}\\exp\\left(-\\frac{n\\pi E_c}{ \\varepsilon}\\right), \\label{probabilityehb} \\end{equation} where the first term $n=1$ corresponds to the Euler-Heisenberg result (\\ref{herate}). Note that the magnetic field produces in the fermionic case an extra factor $\\lfrac{( n\\pi\\beta / \\varepsilon )} {\\tanh ({n\\pi \\beta/ \\varepsilon}})>1$ in each term which enhances the decay rate. The bosonic series (\\ref{probabilityehb}), on the other hand, carries in each term  a suppression factor $\\lfrac{ (n\\pi\\beta / \\varepsilon )} {\\sinh {n\\pi \\beta/ \\varepsilon}}<1$. The average number $\\bar N$ (\\ref{nprobabilitym1}) is given by \\begin{equation} \\bar N=\\sum_{n=1}^\\infty nP_n =V\\Delta t  \\frac{\\alpha}{\\pi} \\frac{ \\alpha\\beta \\varepsilon } {\\tanh {\\pi \\beta/ \\varepsilon}}\\exp\\left(-\\frac{\\pi E_c}{ \\varepsilon}\\right). \\label{nprobabilitym} \\end{equation}  The decay rate $ \\Gamma /V$ gives the number of electron--positron pairs produced per unit volume and time. The prefactor can be estimated on dimensional grounds. It has the dimension of $E_c^2/\\hbar $, i.e., $m^4 c^5/\\hbar^4$. This arises from the energy of a pair $2m_ec^2$ divided by the volume whose diameter is the Compton wavelength $\\hbar/m_ec$, produced within a Compton time $\\hbar/m_ec^2$. The exponential factor suppresses pair production as long as the electric field is much smaller than the critical electric field $E_c$, in  which case the JWKB results (\\ref{wkbehfermion2}) and (\\ref{wkbehboson2}) are good approximations.  The general results ({\\ref{probabilityeh}),(\\ref{probabilityehb}) was first obtained by Schwinger \\cite{1951PhRv...82..664S,1954PhRv...93..615S,1954PhRv...94.1362S} for scalar and spinor electrodynamics (see also Nikishov \\cite{Nikishov1969}, Batalin and Fradkin \\cite{1965JETP...21..375V}). The method was extended to special space-time-dependent fields in Refs.~\\cite{1971ZhPmR..13..261P,1972JETP...34..709P,2001JETPL..74..133P,Narozhnyi:1970uv,1970TMP.....5.1080B}. The monographs \\cite{Itzykson2006,Kleinert2008,Greiner1985,1980ato..book.....G,Fradkin1991} can be consulted about more detailed calculation, discussion and bibliography.  \\subsubsection{Effective nonlinear Lagrangian for arbitrary constant electromagnetic field}\\label{nonlinear}  Starting from the integral form of Heisenberg and Euler Lagrangian (\\ref{oneloopl}) we find explicitly real and imaginary parts of the effective Lagrangian $\\Delta{\\mathcal L}_{\\rm eff}$ (\\ref{oneloopl}) for arbitrary constant electromagnetic fields ${\\bf E}$ and ${\\bf B}$ \\cite{2006JKPSP}. The essential step is to reach a direct analytic form resulting from performing the integration. We use the expressions \\cite{1994tisp.book.....G}, \\begin{align} e\\varepsilon s\\coth(e\\varepsilon s ) &=\\sum_{n=-\\infty}^{\\infty}\\frac{s^2 }{(s^2+\\tau^2_n)};\\quad \\tau_n\\equiv n\\pi/e\\varepsilon,\\label{cosnm0}\\\\ e\\beta s\\cot(e\\beta s ) &=\\sum_{m=-\\infty}^{\\infty}\\frac{s^2 }{(s^2-\\tau^2_m)},\\quad \\tau_m\\equiv m\\pi/e\\beta, \\label{cosnm} \\end{align} and obtain for the finite effective Lagrangian of Heisenberg and Euler integral representation, \\begin{align} \\Delta{\\mathcal L}_{\\rm eff} &\\!=\\! \\frac{1}{ 2(2\\pi)^2}\\sum_{n,m=-\\infty}^{\\infty}{\\!\\!\\!}'\\int_0^\\infty ds\\frac{s}{ \\tau^2_n+\\tau^2_m} \\Big[\\frac{\\bar\\delta_{m0} }{(s^2-\\tau^2_m)}-\\frac{\\bar\\delta_{n0} }{(s^2+\\tau^2_n)}\\Big] e^{-is(m^2_e-i\\eta)}, \\label{onelooplfini} \\end{align} where divergent terms $n\\not=0,m=0$, $n=0,m\\not=0$ and $n=m=0$ are excluded from the sum, as indicated by a prime. The symbol $\\bar\\delta_{ij}\\equiv 1-\\delta_{ij}$ denotes the complimentary Kronecker-$\\delta$ which vanishes for $i=j$. The divergent term with $n=m=0$ is eliminated by the zero-field subtraction in Eq.~(\\ref{oneloopl}), while the divergent terms $n\\not=0,m=0$ and $n=0,m\\not=0$ \\begin{align} \\Delta{\\mathcal L}^{\\rm div}_{\\rm eff} = \\frac{1}{ 2(2\\pi)^2}\\int_0^\\infty \\frac{ds}{ s}e^{-is(m^2_e-i\\eta)}2\\left( \\sum_{m=1}^{\\infty}\\frac{1}{\\tau^2_m}-\\sum_{n=1}^{\\infty} \\frac{1}{\\tau^2_n}\\right), \\label{div} \\end{align} are eliminated by the second subtraction in Eq.~(\\ref{oneloopl}). This can be seen by performing the sums \\begin{equation} \\sum_{n=1}^\\infty\\frac{1}{\\tau^k_n}=\\left(\\frac{e\\varepsilon}{\\pi}\\right)^k\\zeta(k);\\quad \\sum_{n=1}^\\infty\\frac{1}{\\tau^k_m}=\\left(\\frac{e\\beta}{\\pi}\\right)^k\\zeta(k), \\label{sumrule} \\end{equation} where $\\zeta(k)=\\sum_n 1/n^k$ is the Riemann function.  The infinitesimal $i\\eta$ accompanying the mass term in the $s$-integral (\\ref{onelooplfini}) is equivalent to replacing $e^{-is(m^2_e-i\\eta)}$ by $e^{-is(1-i\\eta)m^2_e}$. This implies that $s$ is to be integrated slightly below (above) the real axis for $s>0$ ($s<0$). Equivalently one may shift the $\\tau_m$ ($-\\tau_m$) variables slightly upwards (downwards) to $\\tau_m+i\\eta$ ($-\\tau_m-i\\eta$) in the complex plane.  In order to calculate the finite effective Lagrangian (\\ref{onelooplfini}), the factor $e^{-is(1-i\\eta)m^2_e}$ is divided into its sin and cos parts: \\begin{align} \\Delta{\\mathcal L}^{\\sin}_{\\rm eff} &= \\frac{-i}{ 4(2\\pi)^2}\\sum_{n,m=-\\infty}^{\\infty}{\\!\\!\\!}'\\int_{-\\infty}^\\infty \\frac{sds}{ \\tau^2_n+\\tau^2_m} \\Big[\\frac{\\bar\\delta_{m0} }{(s^2-\\tau^2_m)}-\\frac{\\bar\\delta_{n0} }{(s^2+\\tau^2_n)}\\Big]\\sin[s(1-i\\eta)m^2_e]; \\label{onelooplsin}\\\\ \\Delta{\\mathcal L}^{\\cos}_{\\rm eff} &= \\frac{1}{ 2(2\\pi)^2}\\sum_{n,m=-\\infty}^{\\infty}{\\!\\!\\!}'\\int_0^\\infty \\frac{sds}{ \\tau^2_n+\\tau^2_m} \\Big[\\frac{\\bar\\delta_{m0} }{(s^2-\\tau^2_m)}-\\frac{\\bar\\delta_{n0} }{(s^2+\\tau^2_n)}\\Big] \\cos[s(1-i\\eta)m^2_e]. \\label{onelooplcos0} \\end{align} The sin part (\\ref{onelooplsin}) has an even integrand allowing for an extension of the $s$-integral over the entire $s$-axis. The contours of integration can then be closed by infinite semicircles in the half plane, the integration receives contributions from poles $\\pm \\tau_m, \\pm i\\tau_n$, so that the residue theorem leads to, \\begin{align} \\Delta{\\mathcal L}^{\\sin}_{\\rm eff} &= i\\frac{\\alpha\\varepsilon\\beta }{ 2\\pi} \\sum^\\infty_{n=1}\\frac{1}{n}\\coth\\left(\\frac{n\\pi \\beta}{ \\varepsilon}\\right) \\exp(-n\\pi E_c/\\varepsilon)\\label{sinE}\\\\ &-i\\frac{\\alpha\\varepsilon\\beta }{ 2\\pi}\\sum^\\infty_{m=1}\\frac{1}{m}\\coth\\left(\\frac{m\\pi \\varepsilon}{ \\beta}\\right) \\exp (im\\pi E_c/\\beta)\\label{sinB} \\end{align} The first part (\\ref{sinE}) leads to the exact non-perturbative Schwinger rate (\\ref{probabilityeh}) for pair production.  The second term, as we see below, is canceled by the imaginary part of the cos term. In fact, shifting $s\\rightarrow s-i\\eta$, we rewrite the cos part of effective Lagrangian (\\ref{onelooplcos0}) as \\begin{equation} \\Delta{\\mathcal L}^{\\cos}_{\\rm eff} =\\frac{1}{2(2\\pi)^2} \\sum_{n,m=-\\infty}^{\\infty}{\\!\\!\\!}' \\int_0^\\infty ds  \\frac{\\cos(sm^2_e)}{\\tau^2_n+\\tau^2_m} \\left(\\frac{s\\bar\\delta_{m0} }{ s^2-\\tau^2_m-i\\eta}-\\frac{s\\bar\\delta_{n0} }{ s^2+\\tau^2_n-i\\eta}\\right). \\label{onelooplcos} \\end{equation} In the first term of magnetic part, singularities $s=\\tau_m, (m>0)$ and $s=-\\tau_m, (m<0)$ appear in integrating $s$-axis. We decompose \\begin{align} \\frac{s}{ s^2-\\tau^2_m-i\\eta}=i\\frac{\\pi}{2}\\delta(s-\\tau_m)+i\\frac{\\pi}{2}\\delta(s+\\tau_m) +{\\mathcal P}\\frac{s}{ s^2-\\tau^2_m},\\label{h+} \\end{align} where ${\\mathcal P}$ indicates the principle value under the integral. The integrals over the $\\delta$-functions give \\begin{equation} \\Delta_\\delta{\\mathcal L}^{\\cos}_{\\rm eff}=i\\frac{\\alpha\\varepsilon\\beta }{ 2\\pi}\\sum^\\infty_{m=1}\\frac{1}{m}\\coth\\left(\\frac{m\\pi \\varepsilon}{ \\beta}\\right) \\exp (im\\pi E_c/\\beta), \\label{hagensin} \\end{equation} which exactly cancels the second part (\\ref{sinB}) of the sin part $\\Delta{\\mathcal L}^{\\sin}_{\\rm eff}$.  It remains to find the principle-value integrals in Eq.~(\\ref{onelooplcos}), which corresponds to the real part of the effective Lagrangian \\begin{equation} (\\Delta{\\mathcal L}^{\\cos}_{\\rm eff})_{\\mathcal P} =\\frac{1}{2(2\\pi)^2} \\sum_{n,m=-\\infty}^{\\infty}{\\!\\!\\!}' \\frac{1}{\\tau^2_n+\\tau^2_m}{\\mathcal P}\\int_0^\\infty ds \\cos(sm^2_e) \\left(\\frac{s\\bar\\delta_{m0} }{ s^2-\\tau^2_m}-\\frac{s\\bar\\delta_{n0} }{ s^2+\\tau^2_n}\\right). \\label{onelooplcosp} \\end{equation} We rewrite the cos function as $\\cos(sm^2_e)=(e^{ism^2_e}+e^{-ism^2_e})/2$ and make the rotations of integration contours by $\\pm\\pi/2$ respectively, \\begin{eqnarray} (\\Delta{\\mathcal L}^{\\cos}_{\\rm eff})_{\\mathcal P} &=&\\frac{1}{2(2\\pi)^2} \\sum_{n,m=-\\infty}^{\\infty}{\\!\\!\\!}' \\frac{1}{\\tau^2_n+\\tau^2_m}\\times\\nonumber\\\\ &\\times&{\\mathcal P}\\int_0^\\infty d\\tau \\left(\\frac{\\bar\\delta_{m0}\\tau e^{-\\tau} }{ \\tau^2-(i\\tau_mm_e^2)^2} -\\frac{\\bar\\delta_{n0}\\tau e^{-\\tau} }{ \\tau^2-(\\tau_nm_e^2)^2}\\right). \\label{onelooplcosp1} \\end{eqnarray} Using the formulas (see Secs.~3.354, 8.211.1 and 8.211.2 in Ref.~\\cite{1994tisp.book.....G}) \\begin{equation} J(z) \\equiv {\\mathcal P}\\int^\\infty_0 ds \\frac{s e^{-s}}{ s^2-z^2} = -\\frac{1}{2}\\Big[e^{-z}{\\rm Ei}(z) + e^{z}{\\rm Ei}(-z)\\Big], \\label{J(z)1} \\end{equation} where ${\\rm Ei}(z)$ is the exponential-integral function, \\begin{equation} {\\rm Ei}(z) \\equiv {\\mathcal P}\\int_{-\\infty}^z dt \\frac{e^t}{t}=\\log(-z)+\\sum_{k=1}^\\infty\\frac{z^k}{kk!}, \\label{J(z)2} \\end{equation} we obtain the principal-value integrals (\\ref{onelooplcosp1}), \\begin{equation} (\\Delta{\\mathcal L}^{\\cos}_{\\rm eff})_{\\mathcal P}  = \\frac{1}{2(2\\pi)^2}\\sum_{n,m=-\\infty}^{\\infty}{\\!\\!\\!}' \\frac{1}{ \\tau^2_m+\\tau^2_n} \\Big[\\bar\\delta_{m0}J(i\\tau_m m^2_e)-\\bar\\delta_{n0}J(\\tau_n m^2_e)\\Big]. \\label{pertur} \\end{equation}  Having so obtained the real part of an effective Lagrangian for an arbitrary constant electromagnetic field we recover the usual approximate results by suitable expansion of the exact formula. With the help of the series and asymptotic representation (see formula 8.215 in Ref.~\\cite{1994tisp.book.....G}) of the exponential-integral function ${\\rm Ei}(z)$ for large $z$, corresponding to weak electromagnetic fields ($\\varepsilon\\ll 1, \\beta\\ll 1$), \\begin{equation} J(z) =-\\frac{1}{z^2}-\\frac{6}{z^4}-\\frac{120}{z^6}-\\frac{5040}{z^8}-\\frac{362880}{z^{10}}+\\cdot\\cdot\\cdot , \\label{J(z)3} \\end{equation} and Eq.~(\\ref{pertur}), we find, \\begin{align} (\\Delta{\\mathcal L}^{\\cos}_{\\rm eff})_{\\mathcal P}  &= \\frac{1}{2(2\\pi)^2}\\sum_{n,m=-\\infty}^{\\infty}{\\!\\!\\!}' \\frac{1}{ \\tau^2_m+\\tau^2_n} \\Big\\{\\bar\\delta_{n0}\\left[\\frac{1}{\\tau_n^2m_e^4}+\\frac{6}{\\tau_n^4m_e^8}+\\frac{120}{\\tau_n^6m_e^{12}}+\\cdot\\cdot\\cdot\\right]\\nonumber\\\\ &+\\bar\\delta_{m0}\\left[\\frac{1}{\\tau_m^2m_e^4}-\\frac{6}{\\tau_m^4m_e^8}+\\frac{120}{\\tau_m^6m_e^{12}}+\\cdot\\cdot\\cdot\\right]\\Big\\}. \\label{perexpansion} \\end{align} Applying the summation formulas (\\ref{sumrule}), the weak field expansion (\\ref{perexpansion}) is seen to agree with the Heisenberg and Euler effective Lagrangian \\cite{1936ZPhy...98..714H}, \\begin{align} (\\Delta{\\mathcal L}_{\\rm eff})_{\\mathcal P} &= \\frac{2 \\alpha}{90 \\pi E_c^2}\\left\\{ ({\\bf E}^2\\!-\\!{\\bf B}^2)^2+7 ({\\bf E}\\cdot {\\bf B})^2 \\right\\}\\nonumber\\\\ &+\\frac{2 \\alpha}{315 \\pi^2 E_c^4}\\left\\{ 2({\\bf E}^2\\!-\\!{\\bf B}^2)^3+13({\\bf E}^2\\!-\\!{\\bf B}^2) ({\\bf E}\\cdot {\\bf B})^2 \\right\\}+\\cdot\\cdot\\cdot , \\label{Kleinert1} \\end{align} which is expressed in terms of a powers series of weak electromagnetic fields up to $O(\\alpha^3)$. The expansion coefficients of the terms of order $n$ have the general form $m_e^4/(E_c)^n$. As long as the fields are much smaller than $E_c$, the expansion converges.  On the other hand, we can address the limiting form of the effective Lagrangian (\\ref{pertur}) corresponding to electromagnetic fields ($\\varepsilon\\gg 1, \\beta\\gg 1$). We use the series and asymptotic representation (formulas 8.214.1 and 8.214.2 in Ref.~\\cite{1994tisp.book.....G}) of the exponential-integral function ${\\rm Ei}(z)$ for small $z\\ll 1$, \\begin{equation} J(z) =-\\frac{1}{2}\\Big[e^z\\ln(z)+e^{-z}\\ln(-z)\\Big]+\\gamma_E +{\\mathcal O}(z), \\label{smallz} \\end{equation} with $\\gamma_E=0.577216$ being the Euler-Mascheroni constant, we obtain the leading terms in the strong field expansion of Eq.~(\\ref{pertur}), \\begin{equation} (\\Delta{\\mathcal L}^{\\rm cos}_{\\rm eff})_{\\mathcal P} =\\frac{1}{2(2\\pi)^2}\\sum_{n,m=-\\infty}^{\\infty '} \\frac{1}{ \\tau^2_m+\\tau^2_n}\\Big[\\bar\\delta_{n0}\\ln(\\tau_n m^2_e)-\\bar\\delta_{m0}\\ln(\\tau_m m^2_e)\\Big]+\\cdot\\cdot\\cdot . \\label{strongexp} \\end{equation} In the case of vanishing magnetic field ${\\bf B}=0$ and $m=0$ in Eq.~(\\ref{strongexp}), we have, \\begin{equation} (\\Delta{\\mathcal L}^{\\rm cos}_{\\rm eff})_{\\mathcal P}  =\\frac{1}{2(2\\pi)^2}\\sum_{n=1}^{\\infty} \\frac{1}{ \\tau^2_n} \\ln(\\tau_n m^2_e)+\\cdot\\cdot\\cdot =\\frac{\\alpha E^2}{2\\pi^2}\\sum_{n=1}^{\\infty} \\frac{1}{ n^2} \\ln\\left(\\frac{n\\pi E_c}{E}\\right)+\\cdot\\cdot\\cdot , \\label{strongexpe} \\end{equation} for a strong electric field $\\bf E$. In the case of vanishing electric field ${\\bf E}=0$ and $n=0$ in Eq.~(\\ref{strongexp}), we obtain for the strong magnetic field $\\bf B$, \\begin{equation} (\\Delta{\\mathcal L}^{\\rm cos}_{\\rm eff})_{\\mathcal P}  =-\\frac{1}{2(2\\pi)^2}\\sum_{m=1}^{\\infty} \\frac{1}{ \\tau^2_m} \\ln(\\tau_m m^2_e)+\\cdot\\cdot\\cdot =-\\frac{\\alpha B^2}{2\\pi^2}\\sum_{m=1}^{\\infty} \\frac{1}{ m^2} \\ln\\left(\\frac{m\\pi E_c}{B}\\right)+\\cdot\\cdot\\cdot . \\label{strongexpb} \\end{equation} The ($m=1$) term is the one obtained by Weisskopf \\cite{Weisskopf1936}.  We have presented in Eqs.~(\\ref{sinE}), (\\ref{sinB}), (\\ref{hagensin}), (\\ref{pertur}) closed form results for the one-loop effective Lagrangian $\\Delta{\\mathcal L}_{\\rm eff}$ (\\ref{oneloopl}) for arbitrary strength of constant electromagnetic fields. The results will receive fluctuation corrections from higher loop diagrams. These carry one or more factors $ \\alpha $, $ \\alpha ^2$, \\dots~ and are thus suppressed by factors $1/137$. Thus results are valid for all field strengths with an error no larger than roughly  1\\%. If we include, for example, the two-loop correction, the first term in the Heisenberg and Euler effective Lagrangian (\\ref{Kleinert1}) becomes \\cite{Kleinert2008} % \\begin{equation} (\\Delta{\\mathcal L}_{\\rm eff})_{\\mathcal P} = \\frac{2 \\alpha}{90 \\pi E_c^2}\\left\\{ \\left(1+\\frac{40 \\alpha }{9\\pi}\\right)({\\bf E}^2\\!-\\!{\\bf B}^2)^2+7 \\left(1+\\frac{1315 \\alpha }{252\\pi} \\right)({\\bf E}\\cdot {\\bf B})^2 \\right\\}. \\label{Kleinert2} \\end{equation} % Readers can consult the recent review article \\cite{Dunne:2004nc}, where one finds discussions and computations of the effective Lagrangian at the two-loop level, and \\cite{2000NuPhB.564..591D} for discussion of pair production rate.  \\subsection{Theory of pair production in an alternating field}\\label{alternating}  When the external electromagnetic field $F^{\\rm e}_{\\mu\\nu}$ is space-time-dependent, i.e., $F^{\\rm e}_{\\mu\\nu}=F^{\\rm e}_{\\mu\\nu}({\\bf x},t)$ the exponential in Eq.~(\\ref{probability01}) can no longer be calculated exactly. In  this case, JWKB methods have to be used to calculate pair production rates \\cite{1970PhRvD...2.1191B,1971ZhPmR..13..261P,1972JETP...34..709P,2001JETPL..74..133P,1972JETP...35..659P,1973JETPL..17..368M,Marinov1977,1971ZhPmR..13..261P,1972JETP...34..709P,2001JETPL..74..133P}. The aim of this section is to show how one can use a semi-classical JWKB approach to estimate the rate of pair production in an oscillating electric field as first indicated by Brezin and Itzykson in Ref.~\\cite{1970PhRvD...2.1191B}. They evaluated the production rate of charged boson pairs. The results they obtained can be straightforwardly generalized to charged fermion case, since the spins of charged particles contribute essentially with a counting factor to the final results (see Secs.~\\ref{semi} and \\ref{Schwingerformula}). Thus, let the electromagnetic potential be $\\hat {\\bf z}$ directed, uniform in space and periodic in time with frequency $\\omega_0$: \\begin{equation} A^{\\rm e}_\\mu(x)=(0,0,0,A(t)),\\quad A(t)=\\frac{E}{\\omega_0}\\cos\\omega_0 t. \\label{alterpotential} \\end{equation} Then the electric field is $\\hat {\\bf z}$ directed, uniform in space and periodic in time as well. The electric field strength is given by $E(t)=-\\dot A(t)=E\\sin\\omega_0 t$. It is assumed that the electric field is adiabatically switched on and damped off in a time $T^{\\rm e}$, which is much larger than the period of oscillation $T_0=2\\pi/\\omega_0$. Suppose also that $T_0$ is much larger than the Compton time $2\\pi/\\omega$ of the created particle , i.e., \\begin{equation} T^{\\rm e}\\gg T_0\\gg \\frac{2\\pi}{\\omega}\\simeq \\frac{2\\pi}{m_e}, \\label{thyrachy} \\end{equation} where $\\omega=\\sqrt{|{\\bf p}|^2+m_e^2}$, ${\\bf p}$ being the 3-momentum of the created particle. Furthermore, $eE$ is assumed to be much smaller than $m_e^2$, i.e., $E\\ll E_c$ (see Eq.~(\\ref{wkbc1})).  We have to study the time evolution of a scattered wave function $\\psi(t)$ representing the production of particle and antiparticle pairs in the electromagnetic potential (\\ref{alterpotential}). As usual, an antiparticle can be thought of as a wave-packet moving backward in time. Therefore, for large positive time (forward) only positive energy modes ($\\sim e^{-i\\omega t}$) contribute to $\\psi(t)$. Similarly, for large negative times both positive energy and negative energy modes ($\\sim e^{i\\omega t}$) contribute to $\\psi(t)$ which satisfies the differential equation \\cite{1970PhRvD...2.1191B}: \\begin{equation} \\left[\\frac{d^2}{dt^2}+\\omega^2(t)\\right]\\psi(t)=0, \\label{bscatt} \\end{equation} where the ``variable frequency'' is defined as \\begin{equation} \\omega(t)\\equiv \\left\\{m_e^2+{\\bf p_\\perp}^2+[p_z-eA(t)]^2\\right\\}^{1/2}. \\label{bscattw} \\end{equation} The JWKB method suggests a general solution of the from \\begin{equation} \\psi(t)=\\alpha(t)e^{-i\\chi(t)}+\\beta(t)e^{i\\chi(t)}, \\quad \\chi(t)\\equiv \\int_0^tdt'\\omega(t'), \\label{bscattwave} \\end{equation} where the boundary conditions at large positive and negative times are: \\begin{equation} \\alpha(-\\infty)=1,\\quad \\beta(+\\infty)=0; \\quad \\dot\\chi(\\pm\\infty)=\\omega. \\label{bscattcond} \\end{equation} The backward scattering amplitude $\\beta(t)$ for large negative time ($t\\rightarrow -\\infty$) represents the probability of antiparticle production.  The normalization condition $|\\psi(t)|^2=1$ implies \\begin{equation} \\dot\\alpha(t)e^{-i\\chi(t)}+\\dot\\beta(t)e^{i\\chi(t)}=0. \\end{equation} Eq.~(\\ref{bscatt}) can be written in terms of the scattering amplitudes as \\begin{equation} \\dot\\alpha(t)e^{-i\\chi(t)}-\\dot\\beta(t)e^{i\\chi(t)} =-\\frac{\\dot\\omega(t)}{\\omega(t)}\\left[\\alpha(t)e^{-i\\chi(t)}-\\beta(t)e^{i\\chi(t)}\\right], \\label{bscattwaven} \\end{equation} or, which is the same, \\begin{eqnarray} \\dot\\alpha(t)&=&-\\frac{\\dot\\omega(t)}{2\\omega(t)}\\left[\\alpha(t)-\\beta(t)e^{i2\\chi(t)}\\right], \\label{bscattwaven1}\\\\ \\dot\\beta(t)&=&-\\frac{\\dot\\omega(t)}{2\\omega(t)}\\left[\\beta(t)-\\alpha(t)e^{-i2\\chi(t)}\\right]. \\label{bscattwaven2} \\end{eqnarray} It follows from assumption (\\ref{thyrachy}) that $\\dot\\omega(t)$ vanishes as $|t|\\rightarrow\\infty$, i.e., \\begin{equation} \\frac{\\dot\\omega(t)}{\\omega^2(t)}=\\frac{eE[p_z-eA(t)]}{\\left\\{m_e^2+{\\bf p_\\perp}^2+[p_z-eA(t)]^2\\right\\}^{3/2}} \\ll 1. \\end{equation} More precisely \\begin{equation} \\left|\\frac{\\dot\\omega(t)}{\\omega^2(t)}\\right| <\\frac{eE}{m_e^2+{\\bf p_\\perp}^2}<\\frac{eE}{m_e^2}\\ll 1. \\label{bscattwr} \\end{equation} Therefore, $\\alpha(t)$ and $\\beta(t)$ slowly vary in time and tend to constants as $|t|\\rightarrow\\infty$. The phase $e^{i2\\chi(t)}$ oscillates very rapidly as compared to the variation of $\\alpha(t)$ and $\\beta(t)$, for $\\dot\\chi(t)=\\omega(t)\\gg |\\dot\\omega(t)/\\omega(t)|$. In the zeroth order the oscillating terms in Eqs.~(\\ref{bscattwaven1}), (\\ref{bscattwaven2}) are negligible and one finds \\begin{equation} \\alpha^{(0)}(t)=[\\omega/\\omega(t)]^{1/2}\\simeq 1;\\quad \\beta^{(0)}(t)=0, \\label{1iteration} \\end{equation} which duly satisfy the boundary conditions, and \\begin{equation} \\beta^{(1)}(t)=\\int_t^\\infty dt'\\frac{\\dot\\omega(t')}{2\\omega(t')}e^{-i2\\chi(t')}, \\label{bscattwavesolu} \\end{equation} where (\\ref{thyrachy}) and (\\ref{bscattwr}) have been used. $|\\beta^{(1)}(-\\infty)|^2$ gives information about the probability of particle-antiparticle pair production. Namely, the probability of pair production per unit volume and time is given by \\begin{eqnarray} \\tilde {\\mathcal P} &=& \\lim_{T^{\\rm e}\\rightarrow \\infty}\\frac{1}{T^{\\rm e}} \\int \\frac{d^3k}{(2\\pi)^3}|\\beta^{(1)}(-T^{\\rm e})|^2\\nonumber\\\\ &=&\\int \\frac{d^3k}{(2\\pi)^3}\\lim_{T^{\\rm e}\\rightarrow \\infty}\\frac{1}{T^{\\rm e}} \\left|\\int_{-T^{\\rm e}/2}^{T^{\\rm e}/2} dt'\\frac{\\dot\\omega(t')}{2\\omega(t')}e^{-i2\\chi(t')}\\right|^2. \\label{bscattpro} \\end{eqnarray} Since $\\omega(t)$ is a periodic function with the same frequency $\\omega_0$ as $A(t)$ one can make a Fourier series expansion: \\begin{equation} \\frac{\\dot\\omega(t)}{2\\omega(t)}=\\sum_{n=-\\infty}^{+\\infty}c_n e^{in\\omega_0}. \\label{fouriou} \\end{equation} Defined a renormalized frequency $\\Omega$ via $\\chi(t)=t\\Omega$ one finds \\begin{equation} \\Omega\\equiv \\int^{2\\pi}_0\\frac{dx}{2\\pi}\\left[m_e^2+{\\bf k_\\perp}^2+ \\left(k_3-\\frac{eE}{\\omega_0}\\cos x\\right)^2\\right]^{1/2}. \\label{norw} \\end{equation} so that \\begin{equation} \\lim_{T^{\\rm e}\\rightarrow \\infty}\\frac{1}{T^{\\rm e}} \\left|\\int_{-T^{\\rm e}/2}^{T^{\\rm e}/2} dt'\\frac{\\dot\\omega(t')}{2\\omega(t')}e^{-i2\\chi(t')}\\right|^2 =2\\pi\\sum_n\\delta(n\\omega_0-2\\Omega)|c_n|^2. \\label{bscattdel} \\end{equation} Consequently, the probability of pair production (\\ref{bscattpro}) is, \\begin{equation} \\tilde {\\mathcal P}=\\int \\frac{d^3k}{(2\\pi)^2} \\sum_n\\delta(n\\omega_0-2\\Omega)|c_n|^2=\\int \\frac{d^3k}{(2\\pi)^2\\omega_0} |c_{n^\\circ}|^2, \\label{bscattpro1} \\end{equation} where ${n^\\circ}=2\\Omega/\\omega_0$ and $c_{n^\\circ}$ are determined via Eq.~(\\ref{fouriou}) as \\begin{equation} c_{n^\\circ}=\\int_{-\\pi}^{\\pi} \\frac{dx}{2\\pi} \\frac{\\dot\\omega(x)}{2\\omega(x)}\\exp\\left\\{\\frac{2i}{\\omega_0}\\int^x_0dx' \\left[m_e^2+{\\bf p_\\perp}^2+\\left(p_z-\\frac{eE}{\\omega_0}\\cos(x')\\right)^2\\right]^{1/2}  \\right\\}. \\label{bscattpro2} \\end{equation} The expression for $c_{n^\\circ}$ contains a very rapidly oscillating phase factor with frequency of the order of $m_e/\\omega_0$, and it decreases very rapidly in terms of imaginary time $\\tau=-it$. Its evaluation requires the application of the steepest-descent method in the complex time $x=\\omega_0 t$ plane. This is done by selecting a proper contour turning in a neighborhood of the saddle point and following the steepest-descent line, so as to find the main contributions to the integral in Eq.~(\\ref{bscattpro2}). The saddle point originates from branch points and poles in Eq.~(\\ref{bscattpro2}), which are the zeros of $\\omega(x)$. Mathematical details can be found in Ref.~\\cite{1970PhRvD...2.1191B}. One finds \\begin{equation} \\tilde {\\mathcal P} \\simeq\\frac{\\omega_0}{9}\\int \\frac{d^3k}{(2\\pi)^2} e^{-2A}\\cos^2B, \\label{bscattpro3} \\end{equation} where \\begin{equation} -A+iB=\\frac{2i}{\\omega_0}\\int^{x_0}_0dx' \\left[m_e^2+{\\bf p_\\perp}^2+\\left(p_z-\\frac{eE}{\\omega_0}\\cos(x')\\right)^2\\right]^{1/2}, \\label{complexab} \\end{equation} and the saddle point is $x_0=1/\\pi+i\\sinh^{-1}[(\\omega_0/eE)(m_e^2+{\\bf p_\\perp}^2)^{1/2}]$.  The exponential factor $e^{-2A}$ in Eq.~(\\ref{bscattpro3}) indicates that particle-antiparticle pairs tend to be emitted with small momenta. This allows one to estimate the right-hand side of Eq.~(\\ref{bscattpro3}) as follow: (i) $p_z$ is set equal to zero, moreover, the range of the $p_z$-integration is of the order of $2eE/\\omega_0$ as suggested by the classical equation of motion (\\ref{bscatt}); (ii) $\\cos^2B$ is replaced by its average value $1/2$. As a result, one obtains \\cite{1970PhRvD...2.1191B}, \\begin{equation} \\tilde {\\mathcal P}\\simeq\\frac{(eE)^3}{18\\pi\\omega_0^2}\\int_{\\eta^{-1}}^\\infty du u \\exp\\left[-\\frac{\\pi eE}{\\omega_0^2}u^2g(u)\\right], \\label{bscattpro4} \\end{equation} where $\\eta^{-1}=m_e\\omega_0/(eE)$, $u=(m_e^2+{\\bf p_\\perp}^2)\\omega_0^2/(eE)^2$ and \\begin{equation} g(z) = \\frac{4}{\\pi}\\int_0^1dy \\Big[\\frac{1-y^2}{ 1 + z^{-2} y^2}\\Big]^{1/2}=F\\left(\\frac{1}{2},\\frac{1}{2};2;-z^{-2}\\right), \\label{xlasereta} \\end{equation} where $F(1/2,1/2;2;-z^{-2})={}_2F_1(1/2,1/2;2;-z^{-2})$ is the Gauss hypergeometrical function. The function $u^2g(u)$ is monotonically increasing: \\begin{equation} \\frac{eE}{\\omega_0^2}u^2g(u)\\ge \\frac{eE}{\\omega_0^2}\\eta^{-2}g(\\eta) = \\frac{m_e^2}{eE}g(\\eta)\\gg 1, \\label{intlimit} \\end{equation} which indicates that the integral in (\\ref{bscattpro4}) is strongly dominated by values in a neighborhood of $u=\\eta^{-1}$. This allows one to approximately perform the integration and leads to the rate of pair production of charged bosons \\cite{1970PhRvD...2.1191B}, \\begin{equation} \\tilde {\\mathcal P}_{\\rm boson}\\simeq\\frac{\\alpha E^2}{2\\pi}\\frac{1}{g(\\eta)+\\frac{1}{2\\eta} g'(\\eta)} \\exp\\left[-\\frac{\\pi m_e^2}{eE}g(\\eta)\\right]. \\label{bscattpro5} \\end{equation} Analogously, the rate of pair production of charged fermions can be approximately obtained from Eq.~(\\ref{bscattpro5}) by taking into account two helicity states of fermions (see Secs.~\\ref{semi} and \\ref{Schwingerformula}), \\begin{equation} \\tilde {\\mathcal P}_{\\rm fermion}\\simeq\\frac{\\alpha E^2}{\\pi}\\frac{1}{g(\\eta)+\\frac{1}{2\\eta} g'(\\eta)} \\exp\\left[-\\frac{\\pi m_e^2}{eE}g(\\eta)\\right]. \\label{bscattpro5f} \\end{equation} This formula has played an important role in recent studies of electron and positron pair production by laser beams, which we will discuss in some details in Section~\\ref{Xray}. Momentum spectrum of electrons and positrons, produced from the vacuum, was calculated in \\cite{1971ZhPmR..13..261P,1972JETP...34..709P,1972JETP...35..659P,2001JETPL..74..133P}. For $\\eta\\gg1$ this distribution is concentrated along the direction of electric field, while for $\\eta\\ll1$ it approaches isotropic one.  Unfortunately, it appears very difficult to produce a macroscopic electric field with strength of the order of the critical value (\\ref{critical1}) and lifetime long enough ($\\gg \\hbar/(m_ec^2)$) in any ground laboratory to directly observe the Sauter-Euler-Heisenberg-Schwinger process of electron--positron pair production in vacuum. The same argument applies for the production of any other pair of fermions or bosons. In the following Section, we discuss some ideas to experimentally create a transient electric field $E\\lesssim E_c$ in Earth-bound laboratories, whose lifetime is expected to be long enough (larger than $\\hbar/m_ec^2$) for the pair production process to take place.  \\subsection{Nonlinear Compton scattering and Breit-Wheeler process}  In Section~\\ref{BW}, we have discussed the Breit--Wheeler process \\cite{1934PhRv...46.1087B} in which an electron--positron pair is produced in the collision of two real photons $\\gamma_1 + \\gamma_2 \\rightarrow e^+ +e^-$ (\\ref{2gammaee}). The cross-section they obtained is $O(r_e^2)$, where $r_e$ is the classical electron radius, see Eq.~(\\ref{bwsection3}). This lowest order photon-photon pair production cross-section is so small that it is difficult to observe creation of pairs in the collision of two high-energy photon beams, even if their center of mass energy is larger than the energy-threshold $2m_ec^2= 1.02$ MeV.  In the previous Sections we have seen that in strong electromagnetic fields in lasers the effective nonlinear terms (\\ref{Kleinert1}) become significant and therefore, the interaction needs not to be limited to initial states of two photons \\cite{1962JMP.....3...59R, 1971PhRvL..26.1072R}. A collective state of many interacting laser photons occurs.  We turn now to two important processes \\cite{1996PhRvL..76.3116B, 1997PhRvL..79.1626B} emerging in the interaction of an ultrarelativistic electron beam with a terawatt laser pulse, performed at SLAC \\cite{1996NIMPA.383..309K}, when strong electromagnetic fields are involved. The first process is the nonlinear Compton scattering, in which an ultrarelativistic electron absorbs multiple photons from the laser field, but emits only a single photon via the process \\begin{equation} e+n\\omega \\rightarrow e' + \\gamma , \\label{process1} \\end{equation} where $\\omega$ represents photons from the strong electromagnetic wave of the laser beam (its frequency being $\\omega$), $n$ indicates the number of absorbed photons and $\\gamma$ represents a high-energy emitted photon (see Eq.~(\\ref{process2}) for {\\it cross symmetry}). The theory of this nonlinear Compton effect (\\ref{process1}) is given in Section \\ref{lighttheoryquantum}. The same process (\\ref{process1}) has been expressed by Bamber et al. \\cite{1999PhRvD..60i2004B} in a  semi-classical framework. The second is the nonlinear Breit--Wheeler process \\begin{equation} \\gamma+n\\omega \\rightarrow e^+ +e^- . \\label{process2} \\end{equation} between this very high-energy photon $\\gamma$ and multiple laser photons: the high-energy photon $\\gamma$, created in the first process, propagates through the laser field and interacts with laser photons $n\\omega$ to produce an electron--positron pair \\cite{1997PhRvL..79.1626B}.  In the electric field $E$ of an intense laser beam, an electron oscillates with the frequency $\\omega$ of the laser and its maximum velocity in unit of the speed of light is given by \\begin{equation} v_{\\rm max}\\gamma_{\\rm max}=\\frac{eE}{m\\omega},\\quad\\quad \\gamma_{\\rm max}=1/\\sqrt{1-v_{\\rm max}^2}. \\label{maxv} \\end{equation} In the case of weak electric field, $v_{\\rm max}\\ll 1$ and the nonrelativistic electron emits the {\\it dipole} radiation well described in linear and perturbative QED. On the other hand, in the case of strong electric fields, $v_{\\rm max}\\rightarrow 1$ and the ultrarelativistically oscillating electron emits {\\it multi-pole} radiation. The radiated power is then a nonlinear function of the intensity of the incident laser beam. Using the maximum velocity $v_{\\rm max}$ of oscillating electrons in the electric field of laser beam, one can characterize the effect of nonlinear Compton scattering by the dimensionless parameter \\begin{equation} \\eta =v_{\\rm max}\\gamma_{\\rm max}=\\frac{eE_{\\rm rms}}{ m\\omega }=\\frac{m_ec^2}{ \\omega\\hbar}\\frac{E_{\\rm rms}}{ E_c}, \\label{eta} \\end{equation} where the subscript `rms' means root-mean-square, with respect to the number of interacting laser photons with scattered electron. The parameter $\\eta$ can be expressed as a Lorentz invariant, \\begin{equation} \\eta^2=\\frac{e^2|\\langle A_\\mu A^\\mu \\rangle|}{ m_e^2}, \\label{eta1} \\end{equation} where $A_\\mu$ is the gauge potential of laser wave, $\\partial^\\mu A_\\mu=0$ and the time-average is taken over one period of laser wave, $\\langle A_\\mu \\rangle =0$ and \\begin{equation} \\langle A_\\mu A^\\mu \\rangle=\\langle (A_\\mu  -\\langle A_\\mu \\rangle)^2\\rangle. \\label{aaverage} \\end{equation} Eq.~(\\ref{eta1}) shows that $\\eta^2$ is the intensity parameter of laser fields, and $\\eta$ in (\\ref{eta}) coincides with the parameter $\\eta$ introduced in Eq.~(\\ref{bscattpro4}) for the pair production in an alternating electric field (see Section~\\ref{alternating}).  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  \\subsection{Quantum description of nonlinear Compton effect}\\label{lighttheoryquantum}  In Refs.~\\cite{1962JMP.....3...59R, 1971PhRvL..26.1072R, Nikishov1964, Nikishov1964a, Nikishov1965, 1967JETP...25.1135N, Nikishov1979, Nikishov1965a, Sengupta1952, 1964PhRv..133..705B, Goldman1964, 1964PhL.....8..103G, Eberly1969, 1982els..book.....B}, the quantum theory of the interaction of free electrons with the field of a strong electromagnetic wave has been studied. The application of quantum perturbation theory to such interaction requires not only that the interaction constant $\\alpha$ should be small but also that field should be sufficiently weak. The characteristic quantity in this respect is the dimensionless invariant ratio $\\eta$, see (\\ref{eta1}). The photon emission processes occurring in the interaction of an electron with the field of a strong electromagnetic wave have been discussed in Ref.~\\cite{1982els..book.....B} for any $\\eta$ value. The method used is based on an exact treatment of this interaction, while the interaction of the electron with the newly emitted photons regarded as a perturbation.  Laser beam is considered as a monochromatic plane wave, described by the gauge potential $A_\\mu(\\phi)$ and $\\phi=kx$, where wave vector $k=(\\omega, {\\bf k})$ $(k^2=0)$ (see Eq.~(\\ref{2p2epotential})). The Dirac equation can be exactly solved \\cite{Wolkow1935} for an electron moving in this field of electromagnetic plane wave of an arbitrary polarization and the normalized wave function of the electron with momentum $p$ is given by (c.e.g. \\cite{1982els..book.....B}), \\begin{eqnarray} \\psi_p &=& \\left[1+\\frac{e}{2(kp)}\\not\\!{k}\\not{\\hspace{-5pt}A}\\right]\\frac{u(p)}{\\sqrt{2q_0}}e^{i\\Phi},\\label{eplanes}\\\\ \\Phi &=& -px-\\int^{kx}_0\\left[\\frac{e}{(kp)}(pA)-\\frac{e^2}{2(kp)}A^2\\right]d\\phi, \\label{ephase} \\end{eqnarray} where $u(p)$ is the solution of free Dirac equation $(\\not\\!{p}-m_e)u(p)=0$ and the time-average value of 4-vector, \\begin{equation} q= p -\\frac{e^2\\langle A^2\\rangle}{2(kp)}k , \\label{tqmomen} \\end{equation} is the kinetic momentum operator in the electron state $\\psi_p$ (\\ref{eplanes}) and the ``effective mass'' $m_*$ of the electron in the field is \\begin{equation} q^2=m_*^2, \\quad m_*=m_e\\sqrt{1+\\eta^2 }, \\label{effectivemass} \\end{equation} where $\\eta^2$ is given by (\\ref{eta1}). The electron becomes ``heavy'' in an oscillating electromagnetic field.  The $S$-matrix element for a transition of the electron from the state $\\psi_p$ to the state $\\psi_{p'}$, with emission of a photon having momentum $k'$ and polarization $\\epsilon'$ is given by (c.e.g. \\cite{1982els..book.....B}) \\begin{eqnarray} S_{fi}&=&-ie\\int \\bar\\psi_{p'}(\\gamma \\epsilon'{}^*)\\psi_p\\frac{e^{ik'x}}{\\sqrt{2\\omega'}}d^4x\\label{eplaneS1}\\\\ &=&\\frac{1}{2\\omega'\\cdot 2q_0\\cdot 2q_0'}\\sum_n M^{(n)}_{fi}(2\\pi)^4i\\delta^{(4)}(nk+q-q'-k'), \\label{eplaneS} \\end{eqnarray} where the integrand in Eq.~(\\ref{eplaneS1}) is expanded in Fourier series and expansion coefficients are in terms of Bessel functions $J_n$, the scattering amplitude $M^{(n)}_{fi}$ in Eq.~(\\ref{eplaneS}) is obtained\\footnote{The explicit expression $M^{(n)}_{fi}$ is not given here for its complexity, see for example \\cite{1982els..book.....B}.} by integrating over $x$. Eq.~(\\ref{eplaneS}) shows that $S_{fi}$ is an infinite sum of terms, each corresponds to an energy-momentum conservation law $nk+q=q'+k'$, indicating an electron ($q$) absorbs $n$-photons ($nk$) and emits another photon ($k'$) of frequency \\begin{equation} \\omega'=\\frac{n\\omega}{1+(n\\omega/m_*)(1-\\cos\\theta)}, \\label{gammaenergy} \\end{equation} in the frame of reference where the electron is at rest (${\\bf q}=0,q_0=m_*$), and $\\theta$ is the angle between $\\bf k$ and $\\bf k'$. Given the $n$th term of the $S$-matrix $S_{fi}$ (\\ref{eplaneS}), the differential probability per unit volume and unit time yields, \\begin{equation} d{\\mathcal P}_{e\\gamma}^{(n)}= \\frac{d^3{\\bf k}'d^3{\\bf q}'}{(2\\pi)^6\\cdot 2\\omega'\\cdot 2q_0\\cdot 2q_0'}|M^{(n)}_{fi}|^2(2\\pi)^4i\\delta^{(4)}(nk+q-q'-k'). \\label{eplaneP} \\end{equation} Integrating over the phase space of final states $\\int d^3{\\bf k}'d^3{\\bf q}'$, one obtains the total probability of emission from unit volume in unit time (circular polarization), \\begin{equation} {\\mathcal P}_{e\\gamma}=\\frac{e^2m_e^2}{4q_0} \\sum_{n=1}^\\infty\\int_0^{\\kappa_n}\\frac{d\\kappa}{(1+\\kappa)^2} \\left[-4J^2_n(z)+\\eta^2\\big(2+\\frac{\\kappa^2}{1+\\kappa}\\big) (J^2_{n+1}+J^2_{n-1}-2J^2_n) \\right], \\label{epprobability} \\end{equation} where $\\kappa=(kk')/(kp')$, $\\kappa_n=2n(kp)/m_*^2$ and Bessel functions $J_n(z)$, \\begin{equation} z=2m_e^2\\frac{\\eta}{(1+\\eta^2)^{1/2}} \\left[\\frac{\\kappa}{\\kappa_n} \\left(1-\\frac{\\kappa}{\\kappa_n}\\right)\\right]^{1/2}, \\label{bargument} \\end{equation} for any $\\eta$ value. A systematic investigation of various quantum processes in the field of a strong electromagnetic wave can be found in \\cite{Nikishov1964, Nikishov1964a, Nikishov1965, 1967JETP...25.1135N, Nikishov1979, Nikishov1965a}, in particular photon emission and pair production in the field of a plane wave with various polarizations are discussed.  We now turn to the Breit--Wheeler process for multi-photons (\\ref{process2}). In this process, the pair production is attributed to the interaction of a high-energy photon with many laser photons in the electromagnetic laser wave. Actually, the Breit--Wheeler process for multi-photons, see Eq. (\\ref{process2}), is related to the nonlinear Compton scattering process, see Eq. (\\ref{process1}), by {\\it crossing symmetry}. By replacement $p\\rightarrow -p$ and $k'\\rightarrow -l$ and reverse the common sign of the expression in Eq.~(\\ref{eplaneS}), one obtains the probability of pair production (\\ref{process2}) by a photon $\\gamma$ (momentum $l$) colliding with $n$ laser photons (momentum $k$) per unit volume in unit time (circular polarization) \\cite{Nikishov1964, Nikishov1964a, Nikishov1965, 1967JETP...25.1135N, Nikishov1979, Nikishov1965a}, \\begin{eqnarray} {\\mathcal P}_{\\gamma\\gamma}&=&\\frac{e^2m_e^2}{16l_0}\\sum_{n>n_0}^\\infty\\int_1^{\\upsilon_n} \\frac{d\\upsilon}{\\upsilon^{3/2}(1+\\upsilon)^{1/2}}\\times\\nonumber\\\\ &\\times&\\left[2J^2_n(z)+\\eta^2(2\\upsilon-1)(J^2_{n+1}+J^2_{n-1}-2J^2_n) \\right], \\label{ggprobability} \\end{eqnarray} where $\\upsilon=(kl)^2/4(kq)(kq')$, $\\upsilon_n=n/n_0$, $n_0=2m_*^2/(kl)$ and Bessel functions $J_n(z)$, \\begin{equation} z=4m_e^2\\frac{\\eta (1+\\eta^2)^{1/2}}{(kl)} \\left[\\frac{\\upsilon}{\\upsilon_n}\\left(1-\\frac{\\upsilon}{\\upsilon_n}\\right)\\right]^{1/2}. \\nonumber \\end{equation} In Eq.~(\\ref{ggprobability}), the number $n$ of laser photons must be larger than $n_0$ ($n>n_0$), which is the energy threshold $n_0(kl)=2m_*^2$ for the process (\\ref{process2}) of pair production to occur. ",
        "conclusions": "\\label{remarks}  We have reviewed three fundamental quantum processes which have highlighted some of the greatest effort in theoretical and experimental physics in last seventy years. They all deal with creation and annihilation of electron--positron pairs. We have followed the original path starting from the classical works of Dirac, on the process $e^+e^-\\rightarrow 2\\gamma$, and the inverse process, $2\\gamma\\rightarrow e^+e^-$, considered by Breit--Wheeler. We have then reviewed the $e^+e^-$ pair creation in a critical electric field $E_c=m_e^2c^3/(\\hbar e)$ and the Sauter-Heisenberg-Euler-Schwinger description of this process both in Quantum Mechanics and Quantum Electro-Dynamics. We have also taken this occasion to reconstruct the exciting conceptual developments, initiated by the Sauter work, enlarged by the Born-Infeld nonlinear electrodynamical approach, finally leading to the Euler and Euler-Heisenberg results. We were guided in this reconstruction by the memories of many discussions of one of us (RR) with Werner Heisenberg. We have then reviewed the latest theoretical developments deriving the general formula for pair production rate in electric fields varying in space and in time,  compared with one in a constant electric field approximation originally studied by Schwinger within QED.  We also reviewed recent studies of pair production rates in selected electric fields varying both in space and in time, obtained in the literature using instanton and  JWKB methods. Special attention has been given to review the pair production rate in electric fields alternating periodically in time, early derived by  Brezin, Itzykson and Popov, and the nonlinear Compton effect in the processes of electrons and photons colliding with laser beams, studied by Nikishov and Narozhny. These theoretical results play an essential role in Laboratory experiments to observe the pair production phenomenon using laser technologies.  We then reviewed the different level of verification of these three processes in experiments carried all over the world. We stressed the success of experimental verification of the Dirac process, by far one of the most prolific and best tested process in the field of physics. We also recalled the study of the hadronic branch in addition to the pure electrodynamical branch originally studied by Dirac, made possible by the introduction of $e^+e^-$ storage rings technology. We then turned to the very exciting current situation which sees possibly the Breit--Wheeler formula reaching its first experimental verification. This result is made possible thanks to the current great developments of laser physics. We reviewed as well the somewhat traumatic situation in the last forty years of the heavy-ion collisions in Darmstadt and Brookhaven, yet unsuccessfully attempting to observe the creation of electron--positron pairs. We also reviewed how this vast experimental program was rooted in the theoretical ideas of Zeldovich, Popov, Greiner and their schools.  We have also recalled the novelty in the field of relativistic astrophysics where we are daily observing the phenomenon of Gamma Ray Bursts \\cite{1999PhR...314..575P, 2005RvMP...76.1143P, 2006RPPh...69.2259M, Ruffini2003, 2007AIPC..910...55R}. These bursts of photons occur in energy range keV to MeV, last about one second and come from astrophysical sources located at a cosmological distance \\cite{1997Natur.387..783C,1997Natur.386..686V, 1998Natur.393...35K, 1998Natur.393...41H, 1998Natur.393...43R}. The energy released is up to $\\sim 10^{55}$ ergs, equivalent to all the energy emitted by all the stars of all the galaxies of the entire visible Universe during that second. It is generally agreed that the energetics of these GRB sources is dominated by a dense plasma of electrons, positrons and photons created during the process of gravitational collapse leading to a Black Hole, see e.g. \\cite{1975PhRvL..35..463D,1999A&A...350..334R,2000A&A...359..855R,1998A&A...338L..87P} and references therein. The Sauter-Heisenberg-Euler-Schwinger vacuum polarization process, we have considered in the first part of the report, is a classic theoretical model to study the creation of an electron--positron optically thick plasma. Similarly the Breit--Wheeler and the Dirac processes we have discussed, are essential in describing the further evolution of such an optically thick electron--positron plasma. The GRBs present an unique opportunity to test new unexplored regime of ultrahigh energy physics with Lorentz factor $\\gamma\\sim 100-1000$ and relativistic field theories in the strongest general relativistic domain.  The aim in this report, in addition to describe the above mentioned three basic quantum processes, has been to identify and review three basic relativistic regimes dealing with an optically thin and optically thick electron--positron plasma.  The first topic contains the basic results of the physics of black holes, of their energetics and of the associated process of vacuum polarization. We reviewed the procedures to generalize in a Kerr--Newman geometry the QED treatment of Schwinger and the creation of enormous number of $10^{60}$ electron--positron pairs in such a process.  The second topic is the back reaction of a newly created electron--positron plasma on an overcritical electric field. Again we reviewed the Breit--Wheeler and Dirac processes applied in the wider context of the Vlasov--Boltzmann--Maxwell equations. To discuss the back reaction of electron--positron pair on external electric fields, we reviewed semi-classical and kinetic theories describing the plasma oscillations using respectively the Dirac-Maxwell equations and the Boltzmann--Vlasov equations. We also reviewed the discussions of plasma oscillations damping due to quantum decoherence and collisions, described by respectively the quantum Boltzmann--Vlasov equation and Boltzmann--Vlasov equation with particle collisions terms. We particularly addressed the study of the influence of the collision processes $e^{+}e^{-}\\rightleftarrows \\gamma\\gamma$ on the plasma oscillations in supercritical electric field $E > E_c$.  After $10^{3-4}$ Compton times, the oscillating electric field is damped to its critical value with a large number of photons created. An equipartition of number and energy between electron--positron pairs and photons is reached. For the plasma oscillation with undercritical electric field $E \\lesssim E_c$,  we recalled that electron--positron pairs, created by the vacuum polarization process, move as charged particles in external electric field reaching a maximum Lorentz factor at finite length of oscillations, instead of arbitrary large Lorentz factors, as traditionally assumed. Finally we point out some recent results which differentiate the case $E>E_{c}$  from the one $E<E_{c}$ with respect to the creation of the rest mass of the pair versus its kinetic energy. For $E>E_{c}$ the vacuum polarization process transforms the electromagnetic energy of the field mainly in the rest mass of pairs, with moderate contribution to their kinetic energy. Such phenomena, certainly fundamental on astrophysical scales, may become soon directly testable in the verification of the Breit--Wheeler process tested in laser experiments in the laboratory.  As the third topic we have reviewed the recent progress in the understanding of thermalization process of an optically thick electron--positron-photon plasma. Numerical integration of relativistic Boltzmann equation with collisional integrals for binary and triple interactions is used to follow the time evolution of such a plasma, in the range of energies per particle between 0.1 and 10 MeV, starting from arbitrary nonequilibrium configuration. It is recalled that there exist two types of equilibria in such a plasma: kinetic equilibrium, when all particles are at the same temperature, but have different nonzero chemical potential of photons, and thermal equilibrium, when chemical potentials vanish. The crucial role of direct and inverse binary and triple interactions in reaching thermal equilibrium is emphasized.  In a forthcoming report we will address how the above mentioned three relativistic processes can be applied to a variety of astrophysical systems including neutron stars formation and gravitational collapse, supernovae explosions and GRBs.  This report is dedicated to the progress of theoretical physics in extreme regimes of relativistic field theories which are on the verge of finding their experimental and observational verification in physics and astrophysics. It is then possible from our review and the many references we have given to gain a basic understanding of this new field of research. The three topics which we have reviewed are closely linked to the three quantum processes currently being tested in precision measurement in the laboratories. The experiments in the laboratories and the astrophysical observations cover complementary aspects which may facilitate a deeper and wider understanding of the nuclear and laser physics processes, of heavy-ion collisions as well as neutron stars formation and gravitational collapse, supernovae and GRBs phenomena. We shall return on such an astrophysical and observational topics in a dedicated forthcoming report.  \\vskip1cm \\begin{center} *\\hskip3cm *\\hskip3cm * \\end{center}  \\vskip1cm  We are witnessing in these times some enormous experimental and observational successes which are going to be the natural ground to test some of the theoretical works which we have reviewed in this report. Among the many experimental progresses being done in particle accelerators worldwide we like to give special mention to two outstanding experimental facilities which are expected to give results in the forthcoming years. We refer here to the National Ignition Facility at the Lawrence Livermore National Laboratory to be soon becoming operational, see e.g. \\cite{2009Sci...324..326C} as well as the corresponding facility in France, the Mega Joule project \\cite{2006JPhy4.133..631G}.  In astrophysics these results will be tested in galactic and extragalactic black holes observed in binary X-ray sources, active galactic nuclei, microquasars and in the process of gravitational collapse to a neutron star and also of two neutron stars to a black hole in GRBs. The progress there is equally remarkable. In the last few days after the completion of this report thanks to the tremendous progress in observational technology for the first time a massive hypergiant star has been identified as the progenitor of the supernova SN 2005gl \\cite{Gal-Yam2009}. In parallel the joint success of observations of the flotilla of X-ray observatories and ground-based large telescopes \\cite{Ruffini2009} have allowed to identify the first object ever observed at $z\\approx8$ the GRB090423 \\cite{2009arXiv0905.0001P}.  To follow the progress of this field we are planning a new report which will be directed to the astrophysical nature of the progenitors and the initial physical conditions leading to the process of the gravitational collapse. There the electrodynamical structure of neutron stars, the phenomenon of the supernova explosion as well as theories of Gamma-Ray Bursts (GRBs) will be discussed. Both in the case of neutron stars and the case of black holes there are fundamental issues still to be understood about the process of gravitational collapse especially with the electrodynamical conditions at the onset of the process. The major difficulties appear to be connected with the fact that all fundamental interactions, the gravitational, the electromagnetic, the strong, the weak interactions appear to participate in essential way to this process which appear to be therefore one of the most interesting fundamental process of theoretical physics. Current progress is presented in the following works \\cite{2007IJMPD..16....1R, 2008AIPC..966..147R, 2008arXiv0804.3197R, 2008AIPC.1053..243R, 2008APS..APR8HE093R, 2008APS..APR8HE095R, 2009arXiv0903.3727P, 2009arXiv0903.4095R, 2009AIPC..........P, 2009AIPC..........R, 2009PhRvD........R, 2009PhRvDa.......R}. What is important to recall at this stage is only that both the supernovae and GRBs processes are among the most energetic and transient phenomena ever observed in the Universe: a supernova can reach energy of $\\sim 10^{52}$ ergs (hypernovae) on a time scale of a few months and GRBs can have emission of up to $\\sim 10^{55}$ ergs \\cite{2009Sci...323.1688A} in a time scale as short as of a few seconds. The central role in their description of neutron stars, for supernova as well as of black holes and the electron--positron plasma discussed in this report, for GRBs, are widely recognized. The reason which makes this last research so important can be seen in historical prospective: the Sun has been the arena to understand the thermonuclear evolution of stars \\cite{1968QB464.B46......}, Cyg X-1 has evidenced the gravitational energy role in explaining an astrophysical system \\cite{2003RvMP...75..995G}, the GRBs are promising to prove the existence for the first time of the ``blackholic energy''. These three quantum processes described in our report reveal the basic phenomena in the process of gravitational collapse predicted by Einstein theory of General Relativity \\cite{Ruffini2009}.  \\vspace{0.3in}"
    },
    "0910/0910.0699_arXiv.txt": {
        "abstract": "{Accurate transition probabilities for forbidden lines are important diagnostic parameters for low-density astrophysical plasmas. In this paper we present experimental atomic data for forbidden [\\ion{Fe}{ii}] transitions that are observed as strong features in astrophysical spectra. } {To measure lifetimes for the $3d^6$($^3$G)$4s$\\,a\\,$^4$G$_{11/2}$ and $3d^6$($^3$D)$4s$\\,b\\,$^4$D$_{1/2}$ metastable levels in \\ion{Fe}{ii} and experimental transition probabilities for the forbidden transitions $3d^7$\\,a\\,$^4$F$_{7/2,9/2}$ -- $3d^6$($^3$G)$4s$\\,a\\,$^4$G$_{11/2}$.} {The lifetimes were measured at the ion storage ring facility CRYRING using a laser probing technique. Astrophysical branching fractions were obtained from spectra of Eta Carinae, obtained with the Space Telescope Imaging Spectrograph onboard the {\\it Hubble Space Telescope}. The lifetimes and branching fractions were combined to yield absolute transition probabilities.} {The lifetimes of the a\\,$^4$G$_{11/2}$ and the b\\,$^4$D$_{1/2}$ levels have been measured and have the following values, $\\tau=0.75\\pm0.10$\\,s and $\\tau=0.54\\pm0.03$\\,s respectively. Furthermore, we have determined the transition probabilities for two forbidden transitions of a\\,$^4$F$_{7/2,9/2}$ -- a\\,$^4$G$_{11/2}$ at 4243.97 and 4346.85\\,\\AA . Both the lifetimes and the transition probabilities are compared to calculated values in the literature.} {} ",
        "introduction": "The cosmic abundance of iron is relatively high compared to other iron group elements and the spectrum of singly ionized iron, \\ion{Fe}{ii}, is a significant contributor to the spectral opacity of the sun and hotter stars. The complex energy level structure of \\ion{Fe}{ii} makes the spectrum extremely line rich and it has been studied in great detail with more than 1000 energy levels identified in the literature (\\citeauthor{Johansson09}\\,\\citeyear{Johansson09}).  \\ion{Fe}{ii} lines are observed in the spectra of a wide variety of astronomical objects, and there is a considerable demand for accurate atomic data for this ion. To meet the accurate data requirements of modern astrophysics, a program was initiated to supply the astronomical community with reliable atomic data: The FERRUM-project (\\citeauthor{FERRUMproj}\\,\\citeyear{FERRUMproj}). The aim of this international collaboration is to measure and evaluate astrophysically relevant experimental and theoretical transition data for the iron group elements.  There are 62 metastable levels in \\ion{Fe}{ii}. The parity forbidden lines from some of these levels are observed as prominent features in astrophysical low density plasmas, such as nebulae, \\ion{H}{ii} regions and circumstellar gas clouds. However, metastable levels have radiative lifetimes several orders of magnitude longer than other levels and are thus more affected by collisions. Due to the absence of these lines in laboratory spectra, the majority of forbidden line transition probabilities ($A$-values) available in the literature are from theoretical calculations.  There are only four metastable levels in \\ion{Fe}{ii} with laboratory measured lifetimes. The a\\,$^6$S$_{5/2}$ and b\\,$^4$D$_{7/2}$ levels have been measured by \\citeauthor{Rostohar}\\,(\\citeyear{Rostohar}) using laser probing of a stored ion beam (a laser probing technique, LPT). In addition, the a\\,$^4$G$_{9/2}$ and b\\,$^2$H$_{11/2}$ levels have been measured by \\citeauthor{Hartman}\\,(\\citeyear{Hartman}) using the LPT. There is good agreement between \\citeauthor{Rostohar}\\,(\\citeyear{Rostohar}) and the calculated values of \\citeauthor{Nussbaumer}\\,(\\citeyear{Nussbaumer}) and \\citeauthor{Quinet}\\,(\\citeyear{Quinet}). However, the lifetimes of \\citeauthor{Rostohar}\\,(\\citeyear{Rostohar}) are systematically shorter than the calculations of Garstang (\\citeyear{Garstang}). \\citeauthor{Hartman}\\,(\\citeyear{Hartman}) also combined the lifetimes of a\\,$^6$S$_{5/2}$, b\\,$^4$D$_{7/2}$ and a\\,$^4$G$_{9/2}$ with branching fractions ($BF$s) to determine experimental $A$-values for forbidden transitions. \\citeauthor{Hartman}\\,\\citeyear{Hartman} measured the $BF$s in astrophysical spectra observed in the ejecta of Eta Carinae and presented additional theoretical $A$-value calculations.  We present radiative lifetimes for the a\\,$^4$G$_{11/2}$ and b\\,$^4$D$_{1/2}$ metastable levels in \\ion{Fe}{ii} measured using the LPT at the CRYRING facility. In addition, $BF$s for two forbidden transitions 4243.97 \\AA\\ and 4346.85 \\AA\\ (a$^4$F$_{9/2}$ - a$^4$G$_{11/2}$ and a$^4$F$_{7/2}$ - a$^4$G$_{11/2}$) have been measured in astrophysical spectra observed in the ejecta of Eta Carinae recorded with the {\\it Hubble Space Telescope} ({\\it HST}) Space Telescope Imaging Spectrograph (STIS). The radiative lifetimes have been combined with the $BF$s to yield $A$-values and we provide a comparison with theoretical values in the literature. ",
        "conclusions": "The relativistic Hartree-Fock (HFR) calculations by \\citeauthor{Quinet}\\,(\\citeyear{Quinet}) show the best general agreement with the experimental values presented in this and previous studies. The agreement is good with the exception of the b\\,$^2$H$_{11/2}$ level (\\citeauthor{Rostohar}\\,\\citeyear{Rostohar}) which differs by $5\\sigma$. This deviation has been explained by \\citeauthor{Hartman}\\,\\citeyear{Hartman} as an effect of level mixing between the b\\,$^2$H$_{11/2}$ and the a\\,$^4$G$_{11/2}$ levels. In addition, this mixing has been observed by \\citeauthor{Johansson}\\,(\\citeyear{Johansson}) through unexpected spin forbidden spectral lines, e.g. the b\\,$^2$H$_{11/2}$ - z\\,$^6$F$_{9/2}$ transition at 6269.97 \\AA .  Comparison of the theoretical lifetime values in Table~\\ref{Lifetimes} for a\\,$^4$G$_{11/2}$ and b\\,$^2$H$_{11/2}$ indicate that the lifetime of the a\\,$^4$G$_{11/2}$ level should be approximately one tenth of the lifetime of b\\,$^2$H$_{11/2}$. However, the experimental values reveal that the a\\,$^4$G$_{11/2}$ level is only one fifth of the lifetime for b\\,$^2$H$_{11/2}$ indicating that the mixing effect can be observed in the lifetimes. Calculations are sensitive to level mixing and performing {\\it ab initio} calculations which reproduce the mixing properties of the system is not trivial. This is illustrated in \\citeauthor{Nahar} (\\citeyear{Nahar}), \\citeauthor{Nahar2} (\\citeyear{Nahar2}), \\citeauthor{HibbertFe} (\\citeyear{HibbertFe}), \\citeauthor{Correge} (\\citeyear{Correge}) and \\citeauthor{Pickering_mixing} (\\citeyear{Pickering_mixing}). In particular, \\citeauthor{Pickering_mixing} (\\citeyear{Pickering_mixing}) emphasized that even though the agreement between the theoretical and experimental \\ion{Fe}{ii} $A$-values investigated \\citeauthor{Pickering_mixing} (\\citeyear{Pickering_mixing}) are in general extremely good, probabilities related to spin forbidden transitions that occur due to level mixing may differ by up to an order of magnitude."
    },
    "0910/0910.4590_arXiv.txt": {
        "abstract": "The WMAP haze is an excess in microwave emission coming from the center of the Milky Way galaxy. In the case of synchrotron emission models of the haze, we present tests for the source of radiating high-energy electrons/positrons. We explore several models in the case of a pulsar population or dark matter annihilation as the source. These morphological signatures of these models are small behind the WMAP Galactic mask, but are testable and constrain the source models. We show that detailed measurements of the morphology may distinguish between the pulsar and dark matter interpretations as well as differentiate among different pulsar models and dark matter profile models individually. Specifically, we find that a zero central density Galactic pulsar population model is in tension with the observed WMAP haze. The Planck Observatory's greater sensitivity and expected smaller Galactic mask should potentially provide a robust signature of the WMAP haze as either a pulsar population or the dark matter. ",
        "introduction": "The Wilkinson Microwave Anisotropy Probe (WMAP) has provided a detailed map of the cosmic microwave background (CMB) as well as the foreground from our Galaxy~\\cite{Hinshaw08,Gold08}. The foreground emission from the center of the Galaxy shows an excess of emission in the inner $5-20\\degree$ around the Galactic center. This excess of microwave emission is known as the ``WMAP haze''~\\cite{Finkbeiner04}. Due to the Galactic plane, the current data on the haze only goes down to about 6 degrees from the Galactic center. Two notable features of the haze are the approximate spherical symmetry and the strong angular dependence of the flux. In particular, the flux increases quickly toward the Galactic center.  The size and shape of the haze is currently not well constrained. Using a different foreground model may change the look of the haze slightly. For instance, Ref.~\\cite{Bottino08} used a different synchrotron map to model the WMAP data and found that the haze in that case is smaller in spatial extent and slightly smaller in magnitude. A cleaner foreground subtraction is needed to further constrain the source of the haze. In this paper, we will use the measurement of the haze presented in Ref.~\\cite{Hooper07}.  An explanation for the WMAP haze is a previously unknown source of microwave emission. The frequency dependence of the haze is quite hard, so a hard source like synchrotron radiation is a likely candidate~\\cite{Finkbeiner07,Dobler07,Hooper07}. The magnetic field in the Galactic center tends to be tens of microgauss. Therefore, the synchrotron radiation from highly relativistic electrons and positrons near the Galactic center could be the source of the microwave signal~\\cite{Ferriere09}. Another possibility for the source of the haze was thermal bremsstrahlung from ionized gas~\\cite{Finkbeiner04}. However, the H$\\alpha$ skymap shows no significant increase in the regions where the haze is strongest. At high density and high temperature ($\\sim 10^{5}\\rm K$) such a gas could still explain the haze, but the emission from such regions is constrained to be small, ruling it out as the source of the haze~\\cite{Kurtz94}.  In the synchrotron emission interpretation, calculating the complete propagation of electrons in the interstellar medium requires the full diffusion-loss equation. This includes spatial diffusion of the $e^{+}e^{-}$, reacceleration of the particles due to momentum-space diffusion, energy loss due to several different mechanisms, convection of the particles in the Galaxy, and the particle source. To solve this complete equation, one may use the software GALPROP~\\cite{Strong98,Strong00}. However, the major necessary physical features of the source and synchrotron emission can be sufficiently modeled using a Green's function solution to the diffusion-loss equation, which we employ~\\cite{Baltz04,Baltz98}.  The source of high energy electrons and positrons required for synchrotron emission remains a mystery~\\cite{Zhang08}. One intriguing explanation of the haze is from dark matter annihilations in the Galactic center~\\cite{Finkbeiner07}. The $e^{+}e^{-}$ produced from by the annihilation products move through the Galactic magnetic field, creating the diffuse synchrotron emission~\\cite{Hooper07}. This model matches the haze well, especially the strong angular dependence and approximate spherical symmetry.  It should be noted that Ref.~\\cite{Cumberbatch09} claims that dark matter annihilations cannot explain the haze because it would require too large of a clumpiness boost factor. In this work, we employ similar methods to those used in that paper and do not find that too large of a boost factor is necessary. Ref.~\\cite{Cumberbatch09} averaged the dark matter density over the direction perpendicular to the Galactic plane, which is inaccurate in light of the morphological effects we present below.  Another source for the haze is $e^{+}e^{-}$ coming from the magnetosphere boundary of Galactic pulsars~\\cite{Kaplinghat09}. This type of emission from nearby pulsars can be the source of positrons in the PAMELA results~\\cite{Adriani08,Boulares89,Hooper08,Yuksel08,Profumo08} and also the source of features seen by ATIC~\\cite{Chang08,Yuksel08,Profumo08} and the Fermi Gamma-Ray Space Telescope~\\cite{Abdo09,Grasso09}.  In this paper, we explore models of the WMAP haze in the context of synchrotron radiation from the electron and positron production of Galactic pulsars and dark matter annihilations. We will show tests of this diffuse synchrotron emission which could distinguish between the exotic explanation of dark matter annihilations and the astrophysical pulsar sources. We also show how haze observations can test pulsar and dark matter models separately. These tests can be done with upcoming results from the Planck Observatory, which has been shown, e.g.\\ in Leach, et al.~\\cite{Leach08}, to be expected to have a much smaller Galactic mask than WMAP~\\cite{Planck06,Stolyarov01}.  Recently, Kaplinghat et al.~\\cite{Kaplinghat09} showed that the morphology of the haze is different depending on the source. This work pointed out that for a source with spherical symmetry, the haze should be slightly elliptically stretched along the Galactic plane while for a centrally-peaked pulsar source, the signal lies primarily along the Galactic plane and has less spherical symmetry.  In this work, we examine the morphological structure of the WMAP haze in detail. First, we consider the signal due to pulsars, both from a Gaussian distribution and one which vanishes at the Galactic center. Second, we look at dark matter annihilations models' sensitivity to the dark matter density profile. Lastly, we compare the pulsar and dark matter scenarios to quantitatively distinguish them from one another. ",
        "conclusions": "High-energy electrons/positrons moving in the Galactic magnetic field could be the source of the WMAP haze. Using the diffusion-loss equation, we have taken a set of models of likely sources and calculated the resulting synchrotron signal as a test for the models. The simplified equation can be solved analytically using a Green's function approach, where it is useful for an understanding of the underlying physics. We considered tests of several models of the haze source, namely a peaked Gaussian distribution of pulsars, a distribution of pulsars peaked away from the Galactic center, and two dark matter distributions as possible sources, all consistent with other constraints.  We find that the WMAP haze could be caused by diffuse synchrotron emission due to pulsars, in agreement with Kaplinghat et al.~\\cite{Kaplinghat09}. Moreover, we find that if the haze is caused by pulsars, the angular flux profile from the Galactic center should be peaked more sharply when using a line-of-sight away from the Galactic plane and is flattened significantly when using a line-of-sight closer to the Galactic plane. Also, if the pulsar distribution vanishes in the Galactic center, a brightness about $20\\degree$ out along the Galactic plane should be visible.  With annihilating dark matter, the angular flux profile from Galactic center is largely independent of direction and has rough spherical symmetry. If the signal is found to be strongly spherically symmetric, then this would be an indication that annihilating dark matter could be the true source.  This directional-dependence of the angular flux profile would be a smoking gun for pulsars causing the haze as opposed to the more spherical dark matter explanation. The Planck probe's increased sensitivity, larger number of bands, enhanced models, and expected smaller Galactic mask will test these models~\\cite{Leach08,Planck06,Stolyarov01}. This will be instrumental in determining the source of the WMAP haze as an astrophysical signal or the indirect detection of the dark matter."
    },
    "0910/0910.4559_arXiv.txt": {
        "abstract": "The geometry of a light wavefront evolving in the 3--space associated with a post-Newtonian relativistic spacetime from a flat wavefront is studied numerically by means of the ray tracing method. For a discretization of the bidimensional wavefront the surface fitting technique is used to determine the curvature of this surface at each vertex of the mesh. The relationship between the curvature of a wavefront and the change of the arrival time at different points on the Earth is also numerically discussed. ",
        "introduction": "\\label{intro} The description of the propagation of light in a gravitational field is even today a central problem in the general theory of relativity. The deflection of light rays and time delays of electromagnetic signals due to the presence of a gravitational field are phenomena detectable with current experimental techniques which allow design new tests for general relativity. In this line, Samuel~\\cite{Sam} recently proposed a method for the direct measurement of the curvature of a light wavefront initially flat, curved when light crosses regions where the gravitational field is non vanishing. He found a relationship between the differences of arrival time recorded at four points on the Earth, measured by employing techniques of very long base interferometry and the volume of a parallelepided determined by four points in the curved wavefront surface. This surface is described by means of a polynomial approximation of the eikonal in a Schwarzschild gravitational field. For more complex gravitational models, such as those considered by Klioner and Peip~\\cite{KP}, de Felice {\\it et al.}~\\cite{dF} or Kopeikin and Sch\\\"afer~\\cite{KS} in studies of light propagation in the solar system, the use of numerical methods for the determination of the geometry of the wavefront surface would also be required. An analytical approach to the relativistic modeling of light propagation has also been developed recently by Le Poncin-Lafitte {\\it et al.}~\\cite{Lep} and Teyssandier and Le Poncin-Lafitte~\\cite{Tey}, where they present methods based on Synge's world function and the perturbative series of powers of the Newtonian gravitational constant, to determine the post-Minkowskian expansions of the time transfer functions.  Nowadays there are numerous techniques in computational differential geometry which allow to analyze geometric properties of surfaces embedded in the ordinary Euclidean space. Techniques of this type are widely applied in different areas such as Computational Geometry, Computer Vision or Seismology. In one of these methods, developed in works by Garimella and Swartz~\\cite{GS} and  Cazals and Pouget~\\cite{CP}, the estimation of differential quantities is established using a fitting of the local representation of the surface by means of a height function given by a Taylor polynomial. A survey of methods for the extraction of quadric surfaces from triangular meshes is found in Petitjean~ \\cite{Pet}.  In this work we consider a discretization of the wavefront surface, replacing this surface by a polyhedral whose faces are equilateral triangles. At initial time, the surface is assumed to be flat and far enough from a gravitational source (say, the Sun) and moving towards this source. We study the deformation of the instantaneous polyhedral representing the wavefront when crossing a region in the relativistic 3--space near the Sun due to the bending of light rays by the gravitational field. In this study, we apply the ray tracing method with initial values on the vertices of the triangular mesh to obtain the corresponding discrete surface at each instant of time. Then we apply the techniques given in \\cite{GS} and \\cite{CP} to describe the wavefront as a surface embedded in the Riemannian 3--space of the post-Newtonian formalism of general relativity. For each vertex in the instantaneous mesh we obtain a quadric which represents locally the surface by applying the least-squares method to the immediate neighboring vertex around the considered point which is represented in normal coordinates adapted to the light rays.  The structure of the paper is as follows: In Section 2 we briefly introduce the basic model for the wavefront propagation in the post-Newtonian formalism. In Section 3, we establish a discretized model of the wavefront surface by means of a regular triangulation and we describe the method employed in this work for the study of the curvature of this surface. In Section 4, a numerical estimation of the curvature of the surface is derived using the ray tracing method. For the numerical integration of the light ray equation, we use the {\\tt Taylor} algorithm implemented by Jorba and Zou \\cite{JZ} which is based on the Taylor series method for the integration of ordinary differential equations and which allow the use of high order numerical integrators and arbitrary arithmetic accuracy, as is required to describe the influence of weak gravitational perturbations on the bending of light rays. Finally, in Section 5, we apply the method discussed above to study the effect of the wavefront curvature on the variation of the arrival time of the light at points on the Earth surface, following the model proposed in Samuel's test \\cite{Sam}. The paper concludes by giving another approach to the estimation of the curvature of the wavefront, derived from an approximation of the Wald curvature \\cite{Blu} associated with a quadruple of points in the wavefront. ",
        "conclusions": "The ray tracing numerical method provides a useful tool for the description of spacelike bidimensional wavefronts  within the framework  of the general relativity. We have studied a method, based on techniques of computational geometry, that allows  to estimate the curvature properties of the surface by making a least-squares fitting of the wavefront surface by a quadric surface in the neighborhood of each point of this surface. The computation of the light rays is carried out using an algorithm based on the Taylor method for the solution of differential equations and  employing high arithmetic precision. Further, we have applied a projection at each step of the numerical integration process that allows to guarantee the fulfillment (at machine precision) of the isotropy condition  for the tangent vector to the light ray. We have also studied numerically the dependence of the curvature properties of the wavefront surface on the value of the Eddington parameter $\\gamma$. On the other hand, we have employed a geometric computational approach to the study of the model proposed by Samuel as a new general relativity test, by determining  a numerical approximation of the volume corresponding to a tetrahedron formed by four points on the wavefront that reaches four receiving stations on the Earth surface. Finally, we have obtained an estimation of the Wald curvature for the wavefront in a vicinity of the Earth by using the differences of arrival time recorded at four receiving stations on the Earth."
    },
    "0910/0910.3465_arXiv.txt": {
        "abstract": "We have performed an On-The-Fly (OTF) mapping survey of ${\\rm ^{12}{CO(J=1-0)}}$ emission in 28 Virgo cluster spiral galaxies using the Five College Radio Astronomy Observatory (FCRAO) 14-m telescope. This survey aims to characterize the CO distribution, kinematics, and luminosity of a large sample of galaxies covering the full extents of stellar disks, rather than sampling only the inner disks or the major axis as was done by many previous single dish and interferometric CO surveys. CO emission is detected in 20 galaxies among the 28 Virgo spirals observed. An atlas consisting of global measures, radial measures, and maps, is presented for each detected galaxy. A note summarizing the CO data is also presented along with relevant information from the literature. The CO properties derived from our OTF observations are presented and compared with the results from the FCRAO Extragalactic CO Survey by Young et al. (1995) which utilized position-switching observations along the major axis and a model fitting method. We find that our OTF derived CO properties agree well with the Young et al. results in many cases, but the Young et al. measurements are larger by a factor of 1.4 - 2.4 for seven (out of 18) cases. We will explore further the possible causes for the discrepancy in the analysis paper currently under preparation. ",
        "introduction": "The high galaxy density and the proximity make the Virgo cluster a particularly interesting laboratory for a galaxy evolution study. Its dynamical evolution is still in progress, and evidence for significant environmental effects is ubiquitous \\citep[e.g., ][]{chung07}. The nearness of the Virgo cluster makes it possible to observe its member galaxies with excellent spatial resolution ($1\\arcsec\\ \\sim 90$ pc). It is the first cluster for which significant HI imaging was done \\citep{vgo84}, and various new HI surveys such as ALFALFA \\citep{gio07} and VIVA \\citep{chu07} have recently been conducted.  The Virgo cluster has also been the subject of many recent multi-wavelength surveys, such as in radio continuum by NRAO VLA Sky Survey \\citep[NVSS;][]{con98}, in H$\\alpha$ \\citep{koo04,che06}, in UV by FAUST \\citep{bro97} and Galaxy Evolution Explorer (GALEX) \\citep[e.g.,][]{dal07}, and in IR by Spitzer \\citep{ken08}. High quality optical images are also available from the Hubble Space Telescope (HST)/ACS Virgo Cluster Survey \\citep{cot04} and Sloan Digital Sky Survey (SDSS).  Here, we present the results from a new imaging survey of $J=1\\rightarrow0\\ ^{12}$CO emission in a large sample of Virgo galaxies in order to address the distribution and characteristics of dense molecular gas in these galaxies. It is well established that molecular clouds are the sites of ongoing star formation (\\citet{lar03} and references therein), and carbon monoxide (CO) is the most commonly used tracer of molecular hydrogen (H$_2$), which is the most abundant but invisible component of cold and dense clouds \\citep[e.g.,][]{sol91}. The global H$_2$ content and its distribution in galaxies and a comparison with other gas and stellar components as a function of morphological type, luminosity, and environment are some of the key insights one can derive from CO observations (see \\citet{you91} and references therein).  Existing CO surveys of Virgo cluster galaxies suffer from limited spatial coverage and small sample sizes.  The FCRAO extragalactic survey \\citep{young95} is the first CO survey covering a wide range of distance, diameter, morphological type, and blue luminosity of some 300 galaxies conducted using the FCRAO 14-m telescope, including CO measurements (detections and upper limits) of 65 Virgo galaxies. Because of the large observing time required, these observations were conducted in the position-switching mode, primarily along the optical major axis of the disks, and the global CO line luminosity was derived assuming a model distribution. More recently, high angular resolution CO images have been obtained using interferometric measurements by \\citet{sak99} and \\citet{sof03}. These measurements reveal a detailed molecular gas distribution at 100 pc scales, but they are limited only to the central region of galaxies.  The BIMA SONG (Survey Of Nearby Galaxies; \\citet{hel03}) has also carried out an imaging survey of CO emission in several Virgo spiral galaxies, incorporating  the short spacing data from the NRAO 12-m telescope. The total number of Virgo galaxies imaged by the BIMA SONG is small (six), however.  These earlier CO observations have revealed important insights on the molecular interstellar medium (ISM) in these galaxies. For example, CO emission is often centrally concentrated, in contrast to the centrally deficient HI distributions commonly found in these galaxies.  The molecular gas distribution also shows little evidence for any influence of the gas stripping mechanisms \\citep[e.g.,][]{ken89}. No central CO peak \\citep{hel03} or nuclear molecular rings \\citep[e.g.,][]{ion05} are seen in other cases. The influence of cluster environment on the molecular content is still poorly understood \\citep[e.g.,][]{bos06}.  A distinguishing characteristic of our new CO survey is the complete imaging of $^{\\rm 12}$CO (J=1--0) emission of a large sample (28 galaxies) of Virgo spirals using the On-The-Fly (OTF) mapping mode of the Five College Radio Astronomy Observatory (FCRAO) 14-m telescope. Our map size of $10^{\\prime} \\times 10^{\\prime}$ is larger than the optical diameter $\\rm D_{25}$ and it covers the entire stellar disk of each galaxy. The 45\\arcsec\\ angular resolution of the new CO images is well matched to the existing VLA HI data and is well suited for comparison with other high resolution multi-wavelength data. The specific questions we aim to address are :  1) To what extent is the CO distribution governed by the disk dynamics?  2) Is there a clear phase transition between HI and $\\rm H_{2}$ as a function of the interstellar radiation field?  3) Are there any systematic differences in the CO properties of spiral galaxies in different environments?  We present in this paper the data and the CO atlas from our OTF mapping survey. In section 2, we describe the sample selection. Observations and data reduction process are described in Section 3. The CO atlas and CO properties are presented in Section 4, and our results are compared with those of \\citet{young95}. Molecular gas distribution in individual galaxies is discussed in Section 5, and the summary and conclusion are given in Section 6. The analysis and interpretation of the data addressing the above questions will be presented in our subsequent papers. ",
        "conclusions": ""
    },
    "0910/0910.3186_arXiv.txt": {
        "abstract": "We have analysed the XMM-Newton spectra of SS\\,433 using a standard model of adiabatically and radiatively cooling X-ray jets. The multi-temperature thermal jet model reproduces well the strongest observed emission line fluxes. Fitting the He- and H-like iron line fluxes, we find that the visible blue jet base temperature is $\\approx 17$~keV, the jet kinetic luminosity $L_k \\sim 2 \\cdot 10^{39}$\\,erg/s and the absorbing column density $N_H \\sim 1.5 \\cdot 10^{22}$\\,cm$^{-2}$. All these parameters are in line with the previous studies. The thermal model alone can not reproduce the continuum radiation in the XMM spectral range, the fluorescent iron line and some broad spectral features. Using the thermal jet-plus-reflection model, we find a notable contribution of ionized reflection to the spectrum in the energy range from $\\sim 3$ to 12~keV. The reflecting surface is highly ionized ($\\xi \\sim 300$), the illuminating radiation photon index changes from the flat spectrum ($\\Gamma \\approx 2$) in the 7\\,--\\,12~keV range to $\\Gamma \\approx 1.6$ in the range of 4\\,--\\,7~keV, and to $\\Gamma \\la 1$ in the range of 2\\,--\\,4~keV. We conclude that the reflected spectrum is an evidence of the supercritical disc funnel, where the illuminating radiation comes from deeper funnel regions, to be further reflected in the outer visible funnel walls ($r\\ga 2 \\cdot 10^{11}$ cm). In the multiple scatterings in the funnel, the harder radiation $> 7$~keV may survive absorption, but softer radiation is absorbed, making the illuminating spectrum curved. We have not found any evidences of reflection in the soft 0.8\\,--\\,2~keV energy range, instead, a soft excess is detected, that does not depend on the thermal jet model details. However the soft component spectrum is basically unknown. This soft component might prove to be the direct radiation of the visible funnel wall. It is represented here either as black body radiation with the temperature of $\\theta_{bb} \\approx 0.1$~keV and luminosity of $L_{bb} \\sim 3 \\cdot 10^{37}$\\,erg/s, or with a multicolour funnel (MCF) model. The soft spectral component has about the same parameters as those found in ULXs. ",
        "introduction": "\\label{S:Intro}  SS\\,433 is the only known persistent superaccretor in the Galaxy -- a source of relativistic jets \\citep[][for review]{Fab04}. This is a massive close binary, where the compact star is most probably a black hole \\citep{Giesetal02,Cheretal05,Hill_Gies08,Blundelletal08}. SS\\,433 intrinsic luminosity is estimated to be $\\sim 10^{40}$\\,erg/s, with its maximum located in non-observed UV region. In effect this is a very blue, but heavily absorbed object; its estimated temperature depends on the accretion disc orientation \\citep{Murdin_Clark80,Cheretal82,Dolanetal97}, being $\\sim 50000 - 70000$\\,K when the disc is the most open to the observer. It is important that almost all the observed radiation is formed in the supercritical accretion disc, and the donor star contributes less than 20\\,\\% of the optical radiation. The system's extreme luminosity suggests that the black hole's mass is $\\sim 10$\\,M$_{\\sun}$. Such a big energy budget of the object is supported by a very well measured kinetic luminosity of the jets, $\\sim 10^{39}$\\,erg/s, both in direct X-ray and optical studies of the jets and in the studies of the jet-powered nebula W\\,50 \\citep{Fab04}. At the same time we know that practically all the energy at the accretion onto a relativistic star is released in X-rays. This means that the observed radiation of SS\\,433 was thermalised in the strong wind coming from the supercritical disc.  The wind from a supercritical disc was first predicted by \\citet{ShakSun73} and later confirmed in radiation-hydrodynamic simulations \\citep{Eggumetal88,Okudaetal05,Ohsuga05}. These ideas led to a prediction that SS\\,433, being observed face-on appears as an extremely bright X-ray source, and we may expect an appearance of a new type of X-ray sources in galaxies \\citep{Katz87,FabMesch00,FabMesch01} -- the face-on SS\\,433 stars. It is quite possible that the new type of X-ray sources, ULXs (ultraluminous X-ray sources, \\citet{Roberts07}) are SS\\,433-like objects observed nearly face-on. Their observed X-ray luminosities are $10^{39} - 3 \\cdot 10^{41}$\\,erg/s and they are certainly related to the massive star population. The disc orientation in SS\\,433 is about edge-on, the angle between the disc plane and the line of sight is $10 \\pm 20^{\\circ}$, due to the precessional variations. Therefore we have no chance to observe the funnel in the supercritical disc directly.  The supercritical disc radiation is not isotropic. If one takes into account the geometrical beaming in the disc funnel with the (half) opening angle $\\vartheta_f \\sim 25^{\\circ} - 30^{\\circ}$ \\citep{Ohsuga05} and that a supercritical disc surface radiation is local Eddington \\citep{ShakSun73}, one finds \\citep{Fab_etal06}, that the observed X-ray luminosity of $\\sim 10^{41}$\\,erg/s is expected for the face-on SS433 star. The recent data show \\citep{Stobbart06,Berghea08} that ULXs posses curvature and rather flat X-ray spectra, which are difficult to interpret with a single-component or any other simple model. The supercritical accretion discs are expected to have flat ($\\nu F_{\\nu} \\propto \\nu^0$) X-ray spectrum \\citep{Poutan07}, because the energy release $Q(r)$ in the discs must be $Q(r) \\propto r^{-2}$ ($T(r) \\propto r^{-1/2}$) to make the disc thick due to the radiation pressure.  The X-ray luminosity of SS\\,433 is $\\sim 10^{36}$\\,erg/s (at a distance of 5.5\\,kpc \\citep{Blund_Bowler04}), four orders of magnitude less than the bolometric luminosity. It is believed that all the X-rays come from the cooling X-ray jets. The jet may be easily accelerated by the radiation in the hydrodynamic funnel \\citep[e.g., ][]{Ohsuga05} of the supercritical disc. The jet velocity value, $v_j \\approx 0.26 c$, and its unique stability, where the velocity does not depend on the activity state, indicate that the jet acceleration must be controlled by the line-locking mechanism \\citep{shapiro86,Fab04}. The funnel has to be relatively transparent for the accelerating radiation, as it was confirmed in radiation-hydrodynamic simulatuons \\citep[e.g., ][]{Ohsuga05}.  These reasonings, however, bear a problem, why do not we observe any funnel radiation, missing four orders of magnitude in the X-rays? Even a small amount of gas in the most outer part of the funnel, where we observe the X-ray jet base, may reflect and scatter some part of the direct $\\sim 10^{40}$\\,erg/s of the funnel radiation. The structure of the funnel and the jet acceleration/collimation region is unknown. Apparently, the direct funnel radiation is entirely blocked for the observer. The 'standard' jet model \\citep{brinkmann88,kotani96,marshall02,filippova06} suggests that all the observed X-ray radiation of SS\\,433 is formed in the cooling X-ray jets. However, an analysis of the latest XMM observations \\citep{brinkmann05} led the authors to a conclusion that the 'standard' jet model can not produce the observed X-ray continuum.  In the 'standard' jet model there are two conical anti-parallel jets, considered identical. The jet's gas is optically thin, being in collisional ionization equilibrium and cooling adiabatically (or adiabatically plus radiatively). The jets are observed beginning from the jet base $r_0$ (the distance between the base of the visible jet and the black hole, different for the red and blue jets), which depends on the precessional phase and on the eclipses by the donor star. The gas temperature at the jet base $T_0(r_0)$ is estimated from observations. Precessional and orbital variability of SS\\,433 is well-known \\citep{Fab04}, the jets (and both the accretion disc, and the accretion disc wind) precess with a 164-day period and an amplitude of $\\pm 20^{\\circ}$. At the precessional phase $\\psi \\approx 0$ the disc is the most open to the observer, the angle between the jets and line of sight is $\\sim 60^{\\circ}$ (we refer here to the approximate values because there are nutation-like and sporadic variabilities, both of about $3-5^{\\circ}$). The accretion disc rim (or the opaque  part of the wind) is thick $h/r \\sim1$ \\citep{filippova06}. Comparing the amplitudes of the precessional and orbital variabilities in different X-ray bands, \\citet{Cheretal05} found that both the donor star and the outer disc rim have about the same size, Therefore, one may expect the same radius for the approaching jet base $r_0 \\sim 10^{12}$\\,cm.  In X-ray observations with ASCA and Chandra \\citep{kotani96,marshall02}, dozens of X-ray emission lines formed both in blue and red jets were resolved. The lines are variable both in positions (in accordance with the jet kinematic model) and in intensities, the strongest are He- and H-like $\\text{Fe XXV\\,K}\\alpha$ and $\\text{Fe\\,XXVI\\,L}\\alpha$ iron lines. This allowed the 'iron line diagnostics' of the gas temperature at the jet bases by measuring the line ratio Fe\\,XXV/Fe\\,XXVI. The temperature at the base of the jet was estimated as $\\theta_0 = kT_0 = 10 \\div20$~keV \\citep{kotani96,marshall02,namiki03}. In the latest XMM-Newton observations \\citet{brinkmann05} estimate the temperature of about $\\theta_0 \\sim 17 \\pm 2$~keV, with the main uncertainty coming from the form of the underlying continuum. Fitting the continuum with a thermal bremsstrahlung model \\citep[GINGA data, ][]{brinknann91} gives a notably bigger value, $\\theta_0 \\ga 30 $~keV. The jet base temperature determined both in the line diagnostics and in the continuum fits, drops notably in eclipses, what confirms the cooling jet model. Analysing the numerous ASCA data, \\citet{kotani96} found that the farther part of the receding jet is weakened (probably absorbed in the external gas located in the disc plane) and Nickel is highly overabundant ($\\sim 10$ times) in the jets. The last finding has been confirmed in the XMM observations \\citep{brinkmann05}. Note that some discrepancies may exist as the system SS\\,433 is highly variable and its appearance depends strongly on the precessional phase.  \\citet{brinkmann05} found that the XMM SS\\,433 spectrum can not be fitted with any simple continuum law. The numerous lines point at the thermal model, but a strong continuum curvature indicates the Comptonization. The best formal model of the additional continuum component is a broken power law with a break at $\\sim 7.1$~keV, where the low energy PL rises with energy ($\\Gamma \\sim -1$), the high energy part is rather steep ($\\Gamma \\sim 4$), however, the parameters are not usually well determined in the fits. The overall continuum is too hard to be bremsstrahlung, the electron-electron bremsstrahlung can not fit it even for the temperatures of 40~keV. A highly absorbed additional thermal component gives a too strong $\\sim 7$~kev edge. \\citet{brinkmann05} discuss a model with a single, very broad $\\sim 7$~kev line, which may be formed in the innermost hot jet (otherwise hidden) and Compton down-scattered by colder material. These uncertainties in continuum make problems for the iron-line diagnostics of the jet base temperature.  Practically in all the studies of SS\\,433 X-ray spectra the jet mass loss $\\dot M_j$ and the jet kinetic luminosity $L_k = \\dot M_j v_j^2 /2$ were determined. To estimate $L_k$ one needs to adopt the jet (half) opening angle $\\vartheta_j$. The mass loss rate derived from the X-ray line intensities depends strongly on the opening angle adopted, i.~e. gas density at the jet base, because the emission measure is $\\propto n^2$. The kinetic luminosity values derived from X-ray spectra were quite diverse, up to $10^{42}$\\,erg/s. However, in two latest studies, one from the Chandra data \\citep{marshall02} and the other from the XMM data \\citep{brinkmann05}, the values are $\\sim 3 \\cdot 10^{38}$ and $\\sim 5 \\cdot 10^{39}$\\,erg/s respectively. They are even closer to one another, matching different distances to SS\\,433 adopted by these authors. \\citet{panferov} found $L_k \\sim 10^{39}$\\,erg/s from the jet Balmer line intensities in optical spectra. The mass loss rate in the jets was estimated \\citep{zealey,FabrBor87,dubner} on the base of W\\,50 nebula powered by the jets with approximately the same result, $L_k \\sim 10^{39}$\\,erg/s. Thus the jet kinetic luminosity was measured with relatively good accuracy.  The jet opening angle was directly measured in the studies of X-ray and optical line widths \\citep{Fab04}. \\citet{marshall02} found $\\vartheta_j = 0.61 \\pm 0.03^{\\circ}$ from the Chandra spectra. They assumed that the density is uniform through the jet cone's cross section and $\\vartheta_j$ is the jet's cross sectional radius. \\citet{borfab87} found $\\vartheta_j = 0.82 \\pm 0.14^{\\circ}$ from optical spectra, assuming that the emitting gas density is Gaussian-distributed in the cone's cross section. We give here the values of $\\vartheta_j$ recalculated, using the \\citet{marshall02} definition. Such a coincidence in the X-ray and optical jet opening angles makes the jets pure ballistic from X-ray ($\\sim 10^{12}$\\,cm) to optical ($\\sim 10^{15}$\\,cm) regions. In optical spectra the jet opening angle has been estimated from H$\\alpha$ line profiles during 8-day long observations, the nutation broading has been taken into account and possible jet sporadic activity was diminished. In other Chandra observations \\citep{namiki03,lopez06} the opening angle has been found twice as big and it was greater in the higher temperature (Fe), than in the lower temperature (Si) parts of the jets. Regarding comparatively short time scales of the X-ray observations and probable line broading due to electron scattering \\citep{namiki03} in the hottest parts of the jet, we may adopt $\\vartheta_j \\approx 0.7^{\\circ}$ as a good estimate of the jet opening angle.  In this paper we analyse X-ray spectra of SS\\,433 using XMM data and the 'standard' jet model. We do not try to determine the jet kinetic luminosity, because the X-ray continuum is complex and can not be explained in the 'standard' jet model \\citep{brinkmann05}. Instead we adopt these two main parameters -- $L_k = 10^{39}$\\,erg/s ($\\dot M_j = 3.3 \\cdot 10^{19}$\\,g/s) and $2\\vartheta_j = 1.5^{\\circ}$, as well established and known. We find that the line fluxes can be well matched with the 'standard' jet model and study the additional components of the X-ray spectrum of SS\\,433. ",
        "conclusions": "The goal of the present paper was to find indications of the supercritical accretion disc funnel in SS\\,433 X-ray radiation. The bolometric luminosity of the disc in SS\\,433 is $\\sim 10^{40}$\\,erg/s and all the luminosity must be initially released in X-rays. This is confirmed by the well-established kinetic luminosity of SS\\,433 jets, which is $\\sim 10^{39}$\\,erg/s, and the fact that the jets are formed in the deepest places of the disc funnel close to the black hole. The observed X-ray luminosity of SS\\,433 is $\\sim 10^{36}$\\,erg/s and it is the radiation of the cooling X-ray jets. Both the orientation of SS\\,433 and the visibility conditions do not allow us to see any deep areas of the funnel. However, the strong mass loss of the accretion disc gives us a hope for detecting some indications of the funnel radiation.  We have analysed the XXM spectra of SS\\,433 with a well-known standard model of the adiabatically cooling X-ray jets, taking into account cooling by radiation. We confirm that the thermal jet model reproduces the emission line fluxes quite well. When the disc is most open to the observer, the visible blue jet base is $r_{0} \\approx 2\\cdot10^{11}$\\,cm, the gas temperature at $r_{0}$ is $\\theta_0 \\approx 17$~keV. The IS gas column density is $N_H \\sim 1.5 \\cdot 10^{22}$\\,cm$^{-2}$, which is in good agreement with the IS extinction found in optical and UV observations. We confirm the previous finding \\citep{kotani96,brinkmann05} that the red jet is probably not seen (blocked) in the soft energy range of 0.8~--2.0~keV. However, both the model with blocked the red jet portions, and the model with the whole red jet visible (but with some higher $N_H$) give the same result, that the thermal jet model alone can not explain the soft continuum. We also confirm that Nickel is highly overabundant in the jets \\citep{kotani96,brinkmann05}.  We have adopted the luminosity $L_k = 10^{39}$\\,erg/s and the jet opening angle $2\\vartheta_j = 1.5^{\\circ}$ as well established and known. However, the derived jet visible base $r_{0}$ is too small, it has to be $\\sim 10^{12}$\\,cm to satisfy the observed eclipses by the donor star. Using the scaling formula $r_0 \\propto \\dot M_j^2 / \\Omega_j$ one may find that during these XMM observations the jet kinetic luminosity was $L_k \\sim 2 \\cdot 10^{39}$\\,erg/s. This scaling should not change the thermal jet spectrum.  We find that the thermal jet model alone can not reproduce the continuum radiation in the XMM spectral range. Therefore we use the thermal jet model together with the REFLION ionized reflection model and find quite a good representation of the spectra. Introducing the reflection not only explains the continuum curvature at the energies of $> 3-4$~kev and the fluorescent line of quasi-neutral iron, it reproduces the broad absorption feature at $\\sim 8$~keV recently detected by \\citet{kubota07}. We independently study three energy ranges (2\\,--\\,4, 4\\,--\\,7 and 7\\,--\\,12~keV) with the same thermal jet model to find the reflection parameters. We find that the ionization parameter (\\ref{E:ionpar}) is about the same in these ranges, $\\xi \\sim 300$, which indicates a highly ionized reflection surface. The illuminating radiation photon index changes from flat, $\\Gamma \\approx 2$, in the 7\\,--\\,12~keV range to $\\Gamma \\approx 1.6$ in the range 4\\,--\\,7~keV and to $\\Gamma \\la 1$ in the range 2\\,--\\,4~keV.  We conclude that the additional reflected spectrum is an indication of the funnel radiation. The illuminating radiation spectrum is flat in the range of 7\\,--\\,12~keV. With multiple scatterings in the funnel the hard radiation may survive absorption. The reflected luminosity in this 7\\,--\\,12~keV spectral range is $L_{refl} \\sim 10^{36}$\\,erg/s. The softer ($2-7$~keV) part of the illuminating spectrum (Fig.\\,\\ref{F:add}) carries a trace of absorption. This is assumed to be due to the multiple scatterings in the funnel. It is important to note that an existence of He- and H-like absorption edges at about zero velocity is a mandatory property of the funnel spectrum to be able to produce the observed jet velocity due to the line-locking mechanism.  The supercritical accretion disc spectrum is expected to be flat, however, it extends up to a few keV only. Comptonization may extend and flatten the spectrum to higher energies. The interaction of the high velocity gas moving along the jet axis with the walls' wind makes the conditions for the Comptonization throughout the whole funnel length consistent. Direct evidences of the Comptonized radiation have recently been found in the INTEGRAL observations \\citep{Cheretal05}.  We have not found any evidences of the reflection in the soft 0.8\\,--\\,2.0~keV energy range. The soft excess is detected in the SS\\,433 spectra. We suppose that in the soft X-rays we observe direct radiation of the visible funnel wall. We represented this component as a black body radiation with a temperature of $\\theta_{bb} \\approx 0.1$~keV and luminosity of $L_{BB} \\sim 3 \\cdot 10^{37}$\\,erg/s. The soft excess may be also fitted with a multicolour funnel (MCF) model. Both the soft X-ray luminosity and its derived spectrum strongly depend on the column density $N_H$. The soft component is necessary to add in order to understand the SS\\,433 spectra within the framework of the thermal jet model.  The soft excess is observed in the spectra of ULXs with the soft component temperature of $\\theta \\sim 0.1$~keV, which is identical with that found in SS\\,433. If the ULXs or some of them are nearly face-on versions of SS\\,433, one may adjust their soft X-ray components to the outer funnel walls radiation \\citep{Poutan07}."
    },
    "0910/0910.1183_arXiv.txt": {
        "abstract": "The Exo-Planet Imaging Camera and Spectrograph (EPICS) for the future 42-meter European-Extremely Large Telescope, will enable direct images, and spectra for both young and old Jupiter-mass planets in the infrared. To achieve the required contrast, several coronagraphic concepts -- to remove starlight -- are under investigation: conventional pupil apodization (CPA), apodized-pupil Lyot coronagraph (APLC), dual-zone coronagraph (DZC), four-quadrants phase mask (FQPM), multi-stages FQPM, annular groove phase mask (AGPM), high order optical vortex (OVC), and band-limited coronagraph (BLC). Recent experiment demonstrated the interest of an halftone-dot process -- namely microdots technique -- to generate the adequate transmission profile of pupil apodizers for CPA, APLC, and DZC concepts. Here, we examine the use of this technique to produce band-limited focal plane masks, and present guidelines for the design. Additionally, we present the first near-IR laboratory results with BLCs that confirm the microdots approach as a suitable technique for ground-based observations. ",
        "introduction": "By the end of 2018, challenging projects such as EPICS (Exo-Planet Imaging Camera and Spectrograph, \\citet{EPICS}) for the future 42-meter European-Extremely Large Telescope, or PFI  \\citep[Planet Formation Imager,][]{M06} for the Thirty Meter Telescope (TMT) will enable direct images, and spectra for warm self-luminous and reflected-light Jovian planets. These instruments will operate with eXtreme Adaptive Optics system (XAO) designed for high Strehl ratios (i.e. $\\sim90\\%$ in $H$-band), as for SPHERE \\citep{2008SPIE.7014E..41B} and GPI \\citep{2006SPIE.6272E..18M}, forthcoming planet-finder instruments with first light planned in 2011.  Several coronagraph concepts have been studied extensively in the past years, with the objective of finding optimized designs that can sufficiently suppress the on-axis starlight, allowing faint companions direct detection (e.g. \\citet{Malbet96, SIVA01, Guyon07, Corono}). Among them, the band-limited coronagraph \\citep[][BLC]{2002ApJ...570..900K} has been proposed to completely remove starlight. The BLC has the advantage of being less sensitive to the primary mirror segmentation, unavoidable with ELTs, than for other concepts \\citep{SIVA05}, and to provide achromatic behavior. Additionally, in \\citet{Corono}, we pointed out the interest of such concept on an ELT for either very bright object detection, or for the search of planets at large angular separations.  Several BLC prototypes have been developed during the past years \\citep[e.g.][]{Debes04b, Trauger, Creep06, Trauger2, Moody08} for visible wavelength application. Several technical approaches have been used: (1/) gray-scale pattern written with an high-energy beam sensitive glass (HEBS) using e-beam lithography, (2/) Notch filter pattern -- binary mask -- written with thick Chromium layer on a substrate, dry-etched with high density decoupled plasma, (3/) gray-scale pattern manufactured with vacuum deposited metals and dielectrics. Even if electron-sensitized HEBS glass can accurately accommodate continuous range of transmission, the darkening of the HEBS glass under electron bombardment is accompanied by a determined phase shift, while the technique suffers from a lack of experience in the near-IR. Same constraints apply for the vacuum deposited metal technique. The notch filter has the advantage of being intrinsically achromatic. These designs, consisting of a particular implementation of small structures (stripes of opaque material with width of about some microns) must be finely controlled in size, spacing, and opacity. Mask errors and tolerance are discussed in \\citet{Kuchner03}, where requirements might be strong for near-IR application.  In this paper, we examine the use of a halftone-dot process, namely microdot technique, to reproduce a continuous mask profile as already done for pupil apodizer for SPHERE \\citep{microdots1, microdots2}, and being manufactured for the JWST NIRCam coronagraph \\citep{Krist09}. These masks consist of distributions of opaque square pixels (called dots) to reproduce the continuous transmission of a filter with several advantages: relative ease of manufacture, achromaticity, reproducibility, and ability to generate continuous transmission ranges, without introducing wavefront errors. Besides, mask errors can be easily pre-compensated \\citep{2007JOSAB..24.1268D}.  Section 2 describes the microdot design principle and properties, while Sect. 3 provides guidelines for the design. Section 4 presents monochromatic, and polychromatic results obtained in laboratory in the near-infrared. Finally Sect. 5 concludes on the suitability of the microdots approach for producing BLCs in the context of ground-based instruments. ",
        "conclusions": "We have described the development and laboratory experiment of Band-limited coronagraphs using a microdots design in the near-infrared. In this paper, we provide design guidelines, and demonstrate the microdots technique as a promising solution for BLC for ground-based observations.  We have shown with numerical simulations that although total starlight cancellation is not possible, theoretical contrast offer with the microdot approach are deep enough to guarantee that the BLC will not set a limit on the performance of a ground-based instrument. We note that the theoretical treatment presented in this study does not consider the complexity introduced by potential spectral, and polarization effects of the physical mask.  We identified two sampling configurations suitable for near-IR experiment: $s=16$ and $s=8$, yielding to identical performance in the experiment. Additionally, we pointed out that the interest of the microdots technique in the light of contrast and IWA requirements is not dependent of the function bandwidth ($\\epsilon$) assuming IWA $\\geq$ 3$\\lambda/D$. This already meets EPICS (20-30 mas in H-band), or SPHERE (0.1 arcsec in H-band) standard requirements. However, we note that stronger requirements of the IWA (1 or 2 $\\lambda/D$, ultimate goal of SPHERE), will set a limit on the interest of the technique.  With prototypes we have demonstrated impressive performance, where limitations have been presumably (on the basis of simulations) introduced either by mask error, or wavefront error of the bench. Improvement of the results presented and resumed in Table \\ref{resum} are foreseen -- at least for the peak attenuation -- with a new set of prototypes that will provide a better accuracy of the profile in the outer part of the function (calibration issue of the manufacturing process).  Additionally, these raw polychromatic results presented belong as the first tests of BLC in the near-IR, and were managed without active wavefront correction, nor particular data reduction post-processing. Performance reached in the experiment went beyond to the SPHERE prototype performances, APLC results obtained so far with a microdot apodizer on the same bench, and yield to similar contrast than the ones presented in \\citet{Creep06} (using $4^{th}$ order notch-filter mask, in a monochromatic visible-wavelength domain) although IWA are not exactly identical.  Ultimately, these final prototypes will be implemented on HOT \\citep[the High-Order Testbench, the AO-facility at ESO,][]{HOTbench}, and being compared with others \\citep[FQPM, APLC, and Lyots,][]{HOTcorono, microdots1} with atmospheric turbulence generator and AO-correction, as already initiated with intensive simulations \\citep{Corono}. Results of this experiment will be presented in a forthcoming paper.  Although this study was carried out in the context of R$\\&$D activities for EPICS, it is potentially applicable to upcoming instruments such as SPHERE, or GPI. We note that the interest of the technique presented in the paper for space-based operations, is subject to science cases (IWA and contrast requirements). For instance, high star-planet contrast ($10^{-10}$ in the visible) at less than 0.1$\\arcsec$ (Terrestrial planet), will very likely required a wavefront stability only reachable with a space telescope. In that situation, the technique employed here will impose a limit on the accessible contrast. Notch-filter masks might be more appropriate."
    },
    "0910/0910.4842_arXiv.txt": {
        "abstract": "{ Laser Comb Wavelength calibration  shows that the ThAr one  is locally unreliable with possible deviations of up to 100 \\ms  within one order range, while delivering an overall 1 \\ms accuracy (Wilken et al 2009). Such deviation corresponds to \\daa   $\\approx 7\\cdot 10^{-6}$ for a Fe\\,{\\sc ii}-Mg\\,{\\sc ii} pair. Comparison of line shifts  among the 5 Fe\\,{\\sc ii} lines, with almost identical sensitivity to fine structure constant  changes, offers a  clean way to   directly test  the presence of possible local  wavelength calibration errors of whatever origin.  We analyzed  5 absorption systems, with \\zabs ranging from 1.15 to 2.19 towards 3 bright QSOs. The results show that  while some lines are aligned within 20 \\ms,  others  reveal    large  deviations reaching 200 \\ms or higher  and corresponding to a \\daa  $\\ge 10^{-5}$ level. The origin of these deviations is not clearly  identified but  could be related to the adaptation of wavelength calibration to CCD manufacturing irregularities. These results  suggest that to draw  conclusions from  \\daa analysis based on  one or only few lines must be done with extreme care. ",
        "introduction": "Investigations into possible systematic effects entering  in the measurement of \\daa have been discussed  among others in  Murphy et al (2003), Chand et al (2006), Molaro et al (2008).  Here we focus onto the wavelength calibration error.  The wavelength calibration of the QSO CCD images is made  by comparison with ThAr lamp taken before and after  the QSO frames. The laboratory wavelengths used in the line identification have uncertainties  between 10-150 \\ms (Palmer \\& Engleman 1983, Lovis \\& Pepe 2007 ).  Errors in the wavelength calibration could be  estimated from the mean wavelength-pixel residuals from the polynomial best solution. Residuals vary   from   3-4 m\\AA   (Chand et al 2006) to  1 m\\AA, or even better,     (Levshakov et al 2006, Thompson et al 2009). Murphy et al (2003) treating the ThAr as if they were QSO lines obtained  that the error in the individual absorption systems is typically few times $10^{-6}$ but  there are significant cases up to  $5 \\cdot 10^{-6}$  or higher. When averaged over a  large sample of  simulated systems they obtained  \\daa = $ 0.4 \\pm 0.8\\cdot 10^{-7}$ and concluded that these errors  are not  expected  to drive line shifts in the QSO absorption-line results.  A similar conclusion was  achieved by Chand et al (2006) where  3 and 2 ThAr lines were  taken in proximity of FeII and MgII lines, respectively.  However, the residuals in the wavelength calibration may not reflect  the real wavelength error entirely. Recently an experiment conducted at HARPS with the first Laser Comb  wavelength calibration  in the optical revealed  significant deviations of the ThAr wavelength calibration  with  peak-to-peak deviations of up to 100 \\ms within one echelle order. Similarly in HIRES-Keck observations  a comparison of the ThAr calibration with a I2 self-calibration spectrum revealed significant and not reproducible deviations (Griest et al 2009) ",
        "conclusions": "By using line shifts measurements of the  five Fe\\,{\\sc ii} lines with  identical sensitivity to $\\alpha$ changes we have  studied   possible systematics in the wavelength calibration. We  found that: \\begin{itemize} \\item{In most of the systems one or more lines are deviating   by several hundreds of \\ms. This exceeds the  likely errors estimated from the residuals  and reveals the presence of hidden systematics in the spectral data. } \\item{ These errors could  severely affect the \\daa measure,  in particular  in the methods where \\daa is derived from few lines or where some lines are crucial for the result. To have a reliable measure it is important to have consistent \\daa obtained from several lines for a given absorption system. } \\item{HE 0001-2340, which is one of the QSO used in the present UVES controversy,  reveals very poor quality and the whole sample should be reanalyzed {\\it ab initio} } \\end{itemize}"
    },
    "0910/0910.4904_arXiv.txt": {
        "abstract": "At energies $\\gtrsim$ 2 keV, active galactic nuclei (AGN) are the source of the cosmic X-ray background (CXB).  For AGN population synthesis models to replicate the peak region of the CXB ($\\sim$30 keV), a highly obscured and therefore nearly invisible class of AGN, known as Compton thick (CT) AGN, must be assumed to contribute nearly a third of the CXB.  In order to constrain the CT fraction of AGN and the CT number density we consider several hard X-ray AGN luminosity functions and the contribution of blazars to the CXB.  Following the unified scheme, the radio AGN luminosity function is relativistically beamed to create a radio blazar luminosity function.  An average blazar spectral energy density model is created to transform radio luminosity to X-ray luminosity. We find the blazar contribution to the CXB to be 12$\\%$ in the 0.5-2 keV band, 7.4$\\%$ in the 2-10 keV band, 8.9$\\%$ in the 15-55 keV band, and 100$\\%$ in the MeV region.  When blazars are included in CXB synthesis models, CT AGN are predicted to be roughly one-third of obscured AGN, in contrast to the prediction of one half if blazars are not considered.  Our model implies a BL Lac X-ray duty cycle of $\\sim$13$\\%$, consistent with the concept of intermittent jet activity in low power radio galaxies. ",
        "introduction": "\\label{sect:intro}  Nearly half a century after the cosmic X-ray background (CXB) was discovered \\citep{g62}, the majority of the CXB up to 10 keV has been resolved into distinct point sources by deep observations conducted by {\\em ROSAT}, {\\em Chandra}, and {\\em XMM-Newton} \\citep{h98, m00, g01, g02, h01, a03, w04, bh05, w05}. These discrete sources are active galactic nuclei (AGN), compact extra-galactic sources powered by accretion onto black holes \\citep{lb69,r84}.  As such, the CXB encapsulates the history of accretion onto super massive black holes and provides a powerful tool to aid scientific understanding of accretion processes \\citep{fb92}.  It has been shown that a large portion of this accretion is shrouded from our view by intervening matter along the line of sight \\citep{sw89, cel92, m94, c95, f99, t09}.  For AGN spectral and spatial density models to match the peak of the CXB at $\\sim$30 keV, the models must predict a large number of highly obscured sources known as Compton thick (CT) AGN \\citep{r99, u03, tu05, b06, gch07}, which have a neutral hydrogen column density $N_{\\mathrm{H}}\\gtrsim$ 1.5 $\\times$ 10$^{24}$ cm$^{-2}$ \\citep{t09}, making them practically invisible in the 2-10 keV band \\citep{g94}.  CXB synthesis models predict CT sources make up roughly half of the obscured AGN population \\citep{r99, u03, tu05, b06, gch07}.  In recent years, studies to observationally constrain the CT fraction have been undertaken.  At first, small local studies seemed to agree with the model predictions that half of all obscured AGN are CT \\citep{r99, gua05}.  However, a recent study by \\citet{t09}, using samples from {\\em INTEGRAL} and {\\em Swift} observations and high redshift, IR-selected CT AGN candidates, suggests that CT AGN contribute only about 9$\\%$ of the X-ray background, which contrasts sharply to the prediction by \\citet{gch07} that CT AGN account for nearly a third of the CXB.  \\citet{m09} studied 88 AGN observed by INTEGRAL/IBIS in the 20-40 keV band and found that at least 16$\\%$ of obscured AGN are CT and $\\gtrsim$24$\\%$ of the AGN in their local sample ($z \\leq 0.015$) are CT. If $\\sim$75$\\%$ of local AGN are obscured \\citep{r99, t09}, the local AGN sample of \\citet{m09} suggests that $\\gtrsim$32$\\%$ of obscured AGN are CT.  Given the uncertainity of the CT AGN fraction, smaller classes of CXB contributors must be considered.  The CXB contribution from the small class of AGN known as blazars has previously been ignored by CXB synthesis models and CT AGN fraction predictions, even though blazars are known to emit in a broad range from radio to TeV energies.  To further constrain model predictions of the CT AGN fraction, the blazar class of AGN must be considered.  Blazars are a unique and extreme class of AGN. Unified models of AGN, as summarized by Antonucci (1993) and Urry \\& Padovani (1995), explain blazars as radio galaxies with relativistic jets viewed close to the line of sight.  Flat spectrum radio quasars (FSRQs) are relativistically beamed FRIIs (luminous radio galaxies) and BL Lac objects (BL Lacs) are relativistically beamed FRIs (less luminous radio galaxies).  The features which define blazars (extreme variability, high luminosity, high polarization, and radio core-dominance) are due to the relativistic beaming caused by looking down the relativistic jet of the blazar \\citep{p07b}.  The details of the spectral energy distribution (SED) of blazars is still a topic riddled with uncertainties \\citep{km08}.   The extreme variability that distinguishes the blazar class necessitates simultaneous multi-wavelength observations to understand the spectral properties \\citep{gt08}.  However, the two-hump form of the blazar SED, spanning from radio to $\\gamma$-ray energies, is well known \\citep{u99}.  The lower energy hump is due to synchrotron radiation while the higher energy hump is due to inverse Compton scattering \\citep{up95}.  It has been shown that blazars are significant progenitors of the $\\gamma$-ray background \\citep{g06, nt06, km08}; therefore it is expected that blazars should have a non-negligible contribution to the CXB.  \\citet{g06} predict that blazars should account for 11-12$\\%$ of the soft CXB around 1 keV; however, no estimation is made for the blazar contribution to the peak region of the CXB around 30 keV.  A recent study by \\citet{a09}, based on the three year {\\em Swift}/BAT blazar sample, claims that blazars contribute about 10$\\%$ of the X-ray background in the 2-10 keV band. In the 15-55 keV band \\citet{a09} predict blazars contribute $\\sim$20$\\%$ if blazars are modeled as a single population or $\\sim$9$\\%$ if FSRQs and BL Lacs are modeled as two distinct populations. Both \\citet{g06} and \\citet{a09} found that blazars could contribute 100$\\%$ of the CXB in the MeV band.  Due to uncertainties in the low luminosity end of the AGN hard X-ray luminosity function (HXLF), multiple HXLFs must be considered (e.g., Ueda et al. 2003; La Franca et al. 2005; Silverman et al. 2008; Aird et al. 2009; Ebrero et al. 2009; Yencho et al. 2009) to understand the range of predicted CT AGN.  Recent AGN HXLFs find that luminosity-dependent density evolution (LDDE) provides the best fit to the observational data \\citep{u03, h05, lf05, si08, e09, y09}.  \\citet{aird09} find a new evolutionary model, luminosity and density evolution (LADE), also fits the observational data well. Both the LDDE and LADE models are in keeping with the findings that the scarce, high-luminosity sources, quasars, show sharp positive evolution from $z$ $\\approx$0-2, while less luminous sources, Seyferts, evolve more temperately \\citep{b05, bh05, h05}.  Given the connection between AGN and galaxy evolution (e.g., Ferrarese \\& Merritt 2000; Smol\\v{c}i\\'{c} 2009), it is not surprising that AGN evolution matches the trend of galaxy formation, where massive galaxies formed earlier in cosmological time while smaller structures have waited until more recent times to form (e.g., Cowie et al. 1999).  In this work the blazar contribution to the CXB is predicted and the implications for the CT AGN fraction are discussed in the context of multiple HXLFs. In \\S \\ref{sect:calc} we present the model used for the blazar and non-blazar AGN contributions to the CXB.  In \\S \\ref{sect:res} our results are presented while discussions and conclusions are given in \\S \\ref{sect:sum}.  We assume a $\\Lambda$CDM cosmology with $H_0$ = 70 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\Lambda}$ = 0.7, and $\\Omega_{m}$ = 0.3 \\citep{s07}. ",
        "conclusions": "\\label{sect:sum}  It is clear that blazars make a non-negligible contribution to the CXB and significantly reduce the number of CT AGN predicted, and may be primarily responsible for the MeV background.  This paper presents an upper limit to blazar contribution to the CXB by utilizing the unified model of radio-loud AGN.  The main conclusions found here do not change with a different choice of AGN radio luminosity function (e.g., Condon et al. 2002; Sadler et al. 2002; Best et al. 2005; Kaiser \\& Best 2007; Mauch \\& Sadler 2007); however, beaming parameters need to be modified as these luminosity functions do not distinguish between FRIs and FRIIs or high and low luminosity sources.  A recent study by \\citet{a09}, using the three year sample of {\\em Swift}/BAT blazars, finds similar results for FSRQs as those found here; however, this work finds a greater contribution to the CXB and cosmic $\\gamma$-ray background by BL Lacs.  Due to the small number statistics and small redshift range of the {\\em Swift}/BAT BL Lac sample, \\citet{a09} are not able to uniquely determine the evolutionary parameters and thus assume no evolution for BL Lacs.  This work assumes BL Lacs evolve in the same manner as low luminosity radio galaxies.  Also, \\citet{a09} assume a simple power law SED model for BL Lacs whereas this work utilizes an SED model based on average BL Lac properties taking into account the variety of BL Lac subclasses of LBLs and HBLs.   Several studies have found that BL Lacs contribute substantially to the cosmic $\\gamma$-ray background \\citep{g06, nt06, km08}, thus it is not expected that the BL Lac contribution to the CXB is negligible, as found by \\citet{a09}. As this work uses a more physical BL Lac SED model and a reasonable evolutionary model, we expect that this work may more accurately model the BL Lac contribution to the CXB, although the factor 2.5 discrepancy in the BL Lac source counts in the hard X-ray band must be solved in future works.  \\citet{t09} find the density of CT AGN at $z=0$ with $L_X$ $>$ 10$^{43}$ erg s$^{-1}$ is $\\sim$2.2 $\\times$ 10$^{-6}$ Mpc$^{-3}$.  The luminosity function of \\citet{u03} predicts the density of CT AGN with $L_X$ $>$ 10$^{43}$ erg s$^{-1}$ at $z=0$ to be 7.3 $\\times$ 10$^{-6}$ Mpc$^{-3}$ if blazars are not considered and 4.4 $\\times$ 10$^{-6}$ Mpc$^{-3}$ if blazars are considered.  With the blazar contribution to the CXB considered, the \\citet{u03} over predicts the CT AGN density found by \\citet{t09}, by a factor of 2.  Conversely, the luminosity function proposed by \\citet{e09} predicts the density of CT AGN with $L_X$ $>$ 10$^{43}$ erg s$^{-1}$ at $z=0$ to be 1.1 $\\times$ 10$^{-6}$ Mpc$^{-3}$ if blazars are not considered and 3.6 $\\times$ 10$^{-7}$ Mpc$^{-3}$ if blazars are considered, which is a factor of 6 smaller than the density reported by \\citet{t09}.  According to the INTEGRAL results of \\citet{m09}, the $f_{CT}$ $\\geq$ 0.32 with no upper limit given.  Between different HXLFs there is a large scatter in the predicted $f_{CT}$ and the predicted CT AGN density varies by a factor of ~30.  This clearly illustrates the limits imposed by the uncertainty of the low luminosity end of the AGN HXLF and how important it is for future missions to probe this portion of the HXLF.  It has been shown that blazars, specifically BL Lacs, contribute the majority of the $\\gamma$-ray background \\citep{g06, nt06, km08}.  \\citet{g06} found that unless BL Lacs have a small high energy duty cycle the predicted blazar $\\gamma$-ray emission would over predict the $\\gamma$-ray background. Furthermore, \\citet{km08} suggests that using radio blazar luminosity functions may cause an overestimation of the number of sources emitting robustly at higher energies, as it is not certain that all radio sources will have strong X-ray and $\\gamma$-ray emission. Physical and evolutionary models of quasars indicate that AGN activity is short-lived and possibly recurrent \\citep{s82, cp89, ct92}.  \\citet{fran98} showed that long-lived, continuous AGN activity is not consistent with the black hole mass function they calculated from their sample of 13 local galaxies, but short-lived and recurrent AGN activity matches the data well.  Several sources which appear to be restarted AGNs have been observed \\citep{b83, r93, s99, v04, j07, f09}.  Sources have also been observed which have relic radio lobes but the AGN activity is not currently in an active phase \\citep{parm07, dk09, f09}.  Recent observations by {\\em Hubble Space Telescope} and {\\em Chandra} of the relativistic jet of nearby M87 provide evidence for the intermittent nature of jet X-ray emission (e.g., Perlman et al. 2003; Harris et al. 2006; Stawarz et al. 2006; Madrid 2009).  Large amplitude flaring has been observed from the previously quiescent knot HST-1 in the jet of M87 since 2000 \\citep{mad09}.   This flaring activity is shown to be consistent with shocks occurring within the jet as faster moving particles collide with slower relativistic particles injected into the jet at an earlier time \\citep{perl03, sta06, mad09}.  \\citet{perl03} and \\citet{sta06} suggest the recent X-ray flaring of HST-1 is directly related to material injected at the base of the jet 30-40 years ago.  Therefore, a jet X-ray duty cycle is expected.   Finally, due to the spectral steepening that occurs after the flow of energetic particles into the jet has ceased, the best frequency range to search for relic radio lobes is the low radio regime, less than 1 GHz \\citep{parm07}.  Therefore, it is likely that the radio AGN luminosity function given by \\citet{w01} at 151 MHz includes relic radio lobes. As this would affect the low luminosity end of the luminosity function more prevalently as relic sources tend to not be as luminous as active sources \\citep{dk09}, the BL Lac luminosity function found here may overpredict the number of BL Lacs.  Thus, the average BL Lac X-ray duty cycle is likely to be somewhat larger than the 13$\\%$ found here."
    },
    "0910/0910.0021.txt": {
        "abstract": "We present results from multi-epoch spectral analysis of {\\sl XMM-Newton} and {\\sl Chandra} observations of the broad absorption line (BAL) quasar \\apm. Our analysis shows significant X-ray BALs in all epochs with rest-frame energies lying in the range of $\\sim$ 6.7--18~keV. The X-ray BALs and 0.2--10~keV continuum show significant variability on timescales as short as 3.3~days (proper time) %confirming the intrinsic nature of the absorbing outflow and implying a source size-scale of $\\sim$ 10~$r_{\\rm g}$, where  $r_{\\rm g}$ is the gravitational radius. %The fitted width of the X-ray absorption troughs imply a We find a large gradient in the outflow velocity of the X-ray absorbers with projected outflow velocities of up to 0.76~$c$. %The detected projected maximum outflow velocity of $v \\sim 0.76~c$ The maximum outflow velocity constrains the angle between the wind velocity and our line of sight to be less than $\\sim$ 22$^{\\circ}$. Based on our spectral analysis we identify the following components of the outflow: (a) Highly ionized X-ray absorbing material with an ionization parameter in the range of $2.9 \\simlt \\log\\xi \\simlt 3.9$ (the units of $\\xi$ are erg~cm~s$^{-1}$) and a column density of  $\\log{N_{\\rm H}} \\sim 23$ (the units of  $N_{\\rm H}$ are cm$^{-2}$) outflowing at velocities of up to 0.76~$c$. (b) Low-ionization X-ray absorbing gas with %a column density of $\\log{N_{\\rm H}} \\sim 22.8$. We find a possible trend between the X-ray photon index and the maximum outflow velocity of the ionized %X-ray absorber in the sense that flatter %X-ray spectra appear to result in lower outflow velocities. %We provide a plausible explanation for this effect. Based on our spectral analysis of observations of \\apm\\ over a period of 1.2 years (proper time) we estimate the mass-outflow rate and efficiency of the outflow to have varied between $16_{-8}^{+12}$ $M_{\\odot}$~yr$^{-1}$ and $64_{-40}^{+66}$ $M_{\\odot}$~yr$^{-1}$ and $0.18_{-0.11}^{+0.15}$  to $1.7_{-1.2}^{+1.9}$, respectively. Assuming that the outflow properties of  \\apm\\ are a common property of most quasars at similar redshifts, our results then imply that quasar winds are massive and energetic enough to influence significantly the formation of the host galaxy, provide significant metal enrichment to the interstellar medium (ISM) and intergalactic medium (IGM), and %ISM and IGM, and are a viable mechanism for feedback at redshifts near the peak in the number density of galaxy mergers. ",
        "introduction": "A number of physical mechanisms of interactions of AGNs with their environments have been proposed to explain feedback in galaxies and clusters of galaxies. These feedback mechanisms include radio jets, AGN winds, and AGN radiative heating. %, SN type 1 and starburst winds. It may be the case that several of these mechanisms operate simultaneously and their relative contribution to the feedback process varies depending in part on the size of the super-massive black hole (SMBH), and the combined gravitational potential of the SMBH and host galaxy. The gas and dust content of the host galaxy, the redshift and age of the host, and the inflow and outflow properties of the central super-massive black hole may be important as well.  The radio-jet feedback mechanism was proposed after the discovery of significant radio cavities and shock fronts in several clusters of galaxies. The cavity sizes are large enough to imply that the amount of energy injected in the intracluster medium (ICM), as inferred from the mechanical work $pdV$ required to displace the gas in the radio cavities, is sufficient to balance radiative losses, suppress cooling flows, and halt star formation (e.g., Fabian et al. 2009, and references therein). %Assuming the cavities are inflated by mechanical energy from the radio jets it is found that %the ratio of mechanical energy to radio luminosity ranges by three orders of magnitude. The radio jets are thought to be driven by either accretion onto the black hole (e.g., Blandford \\& Payne 1982) or the black-hole spin (e.g., Blandford \\& Znajek 1977). The details of the process by which radio jets deposit energy and distribute it uniformly throughout the ICM is currently not well understood.  Another possibly important mode that may contribute to feedback is AGN radiation. Radiative heating (e.g., Ciotti \\& Ostriker 2007) has been proposed as an alternative to the radio-jet mechanism to explain the suppression of cooling flows in isolated elliptical galaxies and possibly in clusters of galaxies.  %SN type I and starbust winds are observed in many galaxies and expected to contribute %to the feedback process. %The contribution of SN type I and starburst winds to the feedback process %may be less important in massive galaxies and %clusters of galaxies where the gravitational potential %may be too large for the gas in these relatively low velocity winds to escape. %It is also important to note that most of the observations of SN type I and starbust winds %are of relatively nearby galaxies and little is known for redshifts above $z {\\simgt} 1.5$ where %the number density of mergers is expected to peak.  In field galaxies, especially ones in the redshift range of $z \\approx 1-3 $ where the number density of galaxy mergers is thought to peak (e.g., Di Matteo et al. 2005), quasar winds are thought to be one of the major contributors to feedback. Winds can be driven by radiation pressure, magneto-centrifugal forces, and thermal pressure or a combination of these processes. The potential importance of quasar outflows has been explicitly demonstrated in theoretical models of structure formation and galaxy mergers that incorporate the effects of quasar outflows (e.g., Silk \\& Rees 1998; Scannapieco \\& Oh 2004; Granato et al. 2004; Springel, Di Matteo, \\& Hernquist 2005; Hopkins et al. 2005, 2006). Observational evidence to support these theories of quasar outflows was provided by the discovery of near-relativistic X-ray absorbing winds in several broad absorption line (BAL) quasars (e.g., Chartas et al. 2002, 2003) and later in several non-BAL quasars (e.g., Reeves et al. 2003; Pounds et al. 2003, 2006; Braito et al. 2007).   The relative contribution of each of these proposed feedback mechanisms will vary in part depending on the following conditions:  (a) The importance of the radio-jet mode will depend on the radio loudness and the duty cycle of the central AGN. The presence of a radio jet in an AGN is likely to depend on the available supply of infalling material to feed the black hole and/or the spin of the black hole. We note that several observations imply that the fraction of radio-loud objects varies with redshift and luminosity (e.g., Peacock et al. 1986; Schneider et al. 1992; La Franca et al. 1994; Jiang et al. 2007). Specifically it is found that the radio-loud fraction decreases with redshift and increases with luminosity. The radio jet mode therefore may be more important at later times and in clusters of galaxies and in nearby massive ellipticals that have a larger probability of containing a radio-loud object compared to galaxies at $ z {\\simgt}  1$ which are more likely to contain a radio-quiet AGN.   %(b) A starburst may drive a galactic wind under the condition that the %kinetic energy injected into the disk exceeds its binding energy. %Starburst winds can therefore be an important feedback mechanism %in the case where the escape velocity is low, the %lifetime of the starburst is relatively long and the %mass density of the gas is high.   (b) The feedback process is also expected to depend on the mass of the central black hole and the gravitational potential of the host galaxy. In order for a feedback mechanism to be effective it must be able to drive gas out of the host with a velocity that is larger than the escape velocity determined by the gravitational potential of the system. The mass outflow rates and outflow efficiencies of AGN winds will also depend strongly on the mass and luminosity of the SMBH. Analyses of the kinematic properties of UV absorbers in quasars indicate a strong increase in outflow velocity of UV absorbers with quasar UV luminosity (e.g., Laor \\& Brandt 2002; Ganguly \\& Brotherton 2008). Luminous quasars are found to contain winds with velocities of the UV outflowing material of up to 60,000~km~s$^{-1}$ whereas Seyfert galaxies typically have winds with significantly lower velocities of up to $\\sim$ 2,000 km~s$^{-1}$. %with UV outflowing absorbers of up to ~2000 km/s.   (c) Hydrodynamic simulations of AGN atmospheres driven by radiation pressure indicate that the strengths of the outflows originating from AGN accretion disks depend on the Eddington ratio ($L_{\\rm Bol}/L_{\\rm Edd}$) with more massive and faster winds produced at larger Eddington ratios (e.g., Proga, Stone \\& Drew 1988; Proga, Stone \\& Kallman 2000; Proga \\& Kallman 2004). In systems with a limited supply of gas to fuel the AGN it is expected that AGN winds will be relatively weak.  (d) In the case of a galaxy merger the degree to which each mechanism contributes to the overall feedback process may change with time since the merger event occurred.    Observations in X-rays of the BAL quasar APM~08279+5255, the mini-BAL quasar PG~1115+080, and perhaps the low-ionization BAL quasar H~1413-117 have strongly suggested the presence of near-relativistic outflows of X-ray absorbing material in these objects (Chartas et al. 2002, 2003, 2007a, 2007b). Observations of these quasars in the optical and UV indicate outflow velocities and column densities that are more than an order of magnitude lower than those inferred from the X-ray spectra. These differences in outflow velocities and absorbing column densities suggest that X-ray BALs probe a highly ionized, massive and high-velocity component of the wind that is largely distinct from the absorbers detected in the optical and UV wavebands. Most of the mass and energy of quasar winds is likely carried by the outflowing X-ray absorbers, and observations in the X-ray band are therefore crucial for improving our understanding of the contribution of quasar winds to feedback in galaxies.  This paper will focus on determining the significance of the wind in the $z = 3.91$ gravitationally lensed BAL quasar \\apm\\ in the feedback process during an epoch near the peak of the number density of mergers. Specifically, we present results from analyses of several \\xmm\\ and \\chandra\\ observations of \\apm\\ in order to determine the outflow velocity and its variability, obtain insights on the acceleration mechanism, constrain the geometry of the wind, and estimate the mass outflow rate and the outflow efficiency.  Throughout this paper we adopt a $\\Lambda$-dominated cosmology with $H_{0}$ = 70~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_{\\rm \\Lambda}$ = 0.7, and  $\\Omega_{\\rm M}$ = 0.3. ",
        "conclusions": "We have presented results from an analysis of three \\xmm\\ and two \\chandra\\ observations of the $z = 3.91$ BAL quasar \\apm. The main goals of these observations were to study the kinematic and photoionization properties of the wind in order to assess whether it plays an important role in controlling the evolution of the host galaxy and central black hole. Additional objectives included understanding the mechanism that drives the X-ray absorbers to near-relativistic velocities and constraining the geometry of the outflow.  The main conclusions of our spectral and timing analyses are the following:  (a) X-ray BALs lying in the range of 1.5--3.6~keV (observed frame) are detected at the greater than 99.9\\% confidence level in all five observations of \\apm. We note that a recent analysis of three {\\sl Suzaku} observations of \\apm\\ also shows significant detections of X-ray BALs in all three observations (Saez et al. 2009).   (b) Based on our spectral analysis we identify the following components of the outflow. First, a highly ionized X-ray absorbing material with an ionization parameter in the range of $2.9 \\simlt \\log\\xi \\simlt 3.9$ and a column density of  $\\log{N_{\\rm H}} \\sim 23$ outflowing at velocities of up to 0.76~$c$, and second a low-ionization X-ray absorbing gas with a column density of $\\log{N_{\\rm H}} \\sim 22.8$  that may constitutes the X-ray shielding gas that is thought to protect the UV absorber from becoming over-ionized. Models that include partial covering result in even larger column densities of the ionized absorber of up to $\\log{N_{\\rm H}} \\sim 24$, however, we caution that these values are not well constrained (see model 9 of Table~3) by the spectral fits. We note that no iron overabundance is required to fit the X-ray spectra of \\apm\\ consistent with the spectral analysis of previous observations of this object.   (c) Significant variability of the X-ray continuum and X-ray BALs over short and long time-scales is detected. Specifically, we detect a  $\\sim$ 36.8 $\\pm$ 0.3 \\% change of the 0.2--10~keV pn count-rate between epochs 3 and 4 that are separated by only 3.3~days (proper-time). Based on a simple light-travel time argument this variability time-scale implies a radius of the emission region of $\\sim$ 7.4 $\\times$ 10$^{15}$~cm which is comparable to $r_{\\rm ISCO}$ = 4.5 $\\times$ 10$^{15}$~cm for the case of a Schwarzschild black hole in \\apm. We speculate based on the energy independent spectral change between epochs 3 and 4 (with the exception of the Fe BAL region) that the cause of this variability is a change in the covering fraction as the outflowing absorbing material moves across our line of sight. Significant changes in the flux densities between epochs 2 and 3 that are separated by $\\sim$ 5.4 years (observed-frame) are detected. The variation of flux density is significant below 3~keV (observed frame) showing a increase toward lower energies but is near zero above 3~keV. This type of energy-dependent change suggests that the intrinsic unabsorbed luminosity of \\apm\\ has not significantly varied between these epochs and the long-term variability is the result of a decrease in the opacity of the absorber. This interpretation is consistent with the results of our spectral analysis.   (d) Assuming our interpretation that the high-energy X-ray BALs are produced by resonance transitions of highly-ionized iron %Fe XXV(1s--2p) we estimate the mass-outflow rate and efficiency of the outflow to have varied over a period of 1.2 years (proper-time) between $16_{-8}^{+12}$ $M_{\\odot}$~yr$^{-1}$ and $64_{-40}^{+66}$ $M_{\\odot}$~yr$^{-1}$ and $0.18_{-0.11}^{+0.15}$  to $1.7_{-1.2}^{+1.9}$, respectively.    (e) The detected projected maximum outflow velocity of $v_{\\rm max} \\sim 0.76c$ leads to a constraint on the angle between the outflow direction and the observed line of sight through the absorber of $\\theta_{\\rm max}$ $\\simlt$ 22$^{\\circ}$ (assuming \\FeXXV($1s^2-1s2p$)) and $\\theta_{\\rm max}$ $\\simlt$ 26$^{\\circ}$ (assuming \\FeXXVI($1s-2p$)). Such a small angle is consistent with the unification scheme of BAL and non-BAL quasars that posits that BAL quasars are viewed almost along the outflow direction.   (f) A possible trend is found between the X-ray photon index and the maximum outflow velocity of the ionized X-ray absorber in the sense that flatter X-ray spectra appear to result in lower outflow velocities. One possible explanation is that flatter X-ray spectra over-ionize the X-ray absorber resulting in a decrease of the force multiplier of the X-ray absorber and thus a lower outflow velocity."
    },
    "0910/0910.2305_arXiv.txt": {
        "abstract": "We present FLAMES/GIRAFFE spectroscopy obtained at the Very Large Telescope (VLT). Using these observations we have been able for the first time to observe the Li I doublet in the Main Sequence stars of a Globular Cluster. We also observed Li in a sample of Sub-Giant stars of the same B-V colour.  Our final sample is composed of 84 SG stars and 79 MS stars. In spite of the fact that SG and MS span the same temperature range we find that the equivalent widths of the Li I doublet in SG stars are systematically larger than those in MS stars, suggesting a higher Li content among SG stars. This is confirmed by our quantitative analysis which makes use of both 1D and 3D model atmospheres.  We find that SG stars show, on average, a Li abundance higher by 0.1\\,dex than MS stars. We also detect a positive slope of Li abundance with effective temperature, the higher the temperature the higher the Li abundance, both for SG and MS stars, although the slope is slightly steeper for MS stars. These results provide an unambiguous evidence that the Li abundance changes with evolutionary status.  The physical mechanisms that contribute to this are not yet clear, since none of the proposed models seems to describe accurately the observations. Whether such mechanism can explain the cosmological lithium problem, is still an open question. ",
        "introduction": "The primordial Li abundance inferred from the fluctuations of cosmic microwave background measured by the WMAP satellite (\\cite[Spergel et al. 2007, Cyburt et al. 2008]{spe07,cyburt}) is $\\log ({\\rm Li}/{\\rm H})+ 12 =2.72\\pm0.06$, approximately 0.3--0.5 dex higher than the Li abundance determined in metal-poor stars of the Galactic halo.  Many studies have tried to explain this difference: (a) \\cite[Piau et al. 2006]{pia06} propose that the first generation of stars, Population III stars, could have processed some fraction of the halo gas, lowering the lithium abundance; (b) other authors suggest that the primordial Li abundance has been uniformly depleted in the atmospheres of metal-poor dwarfs by some physical mechanism (e.g. turbulent diffusion as in \\cite{ric05,kor06}; gravitational waves as in \\cite{cat05}, etc.); and (c) finally, it has been also suggested that the standard Big Bang nucleosynthesis (SBBN) calculations should be revised, possibly with the introduction of new physics as in e.g. \\cite{jed04,jed06,jittoh,hisano}.  Here we present the Li abundances of subgiant (SG) and main-sequence (MS) stars of the cluster NGC~6397. This study represents the first Li abundance measurements in MS stars of a Globular Cluster. ",
        "conclusions": "In Fig.~\\ref{figali} we display the derived 1D-NLTE Li abundances versus the 1D effective temperatures of dwarf and subgiant stars of the globular cluster NGC 6397 (see Fig.~2 in \\cite{gon09} for a similar picture but with $T_{\\rm eff}$ and Li abundances computed using 3D models). The stars have been divided into five bins of effective temperature. The error bar in the Li abundance is the dispersion for each bin, divided by the square root of the number of stars in the bin. The Li abundance decreases with decreasing temperature. This lithium abundance pattern is different from what is found among field stars (\\cite[Mel\\'endez \\& Ram{\\'\\i}rez 2004]{mel04}, \\cite[Bonifacio et al. 2007]{bon07}, \\cite[Gonz\\'alez Hern\\'andez et al. 2008]{gon08}).  The difference in the Li abundance of dwarfs and subgiants is $\\sim0.14$~dex if one computes the mean 1D A(Li) and the standard deviation of the mean for the two samples. For subgiants we find mean 1D Li abundances of $2.31\\pm0.01$, while for dwarfs $2.17\\pm0.01$. The 3D models provide similar results although slightly higher Li abundances. \\cite{lin09} also find such difference between the mean Li abundances in MSs and SGs, but only by 0.03~dex although still significant at 1$\\sigma$. However, this result is partly affected by the very narrow range of $T_{\\rm eff}$ for MSs deduced by \\cite{lin09} ($\\sim 80$~K) compared to the wide range ($\\sim 450$~K) for the SGs (see Fig.~7 online in \\cite[Gonz\\'alez Hern\\'andez et al. 2009]{gon09}).  Our results imply that the Li surface abundance depends on the evolutionary status of the star. In Fig.~\\ref{figali} we show the Li isochrones for different turbulent diffusion models (\\cite[Richard et al. 2005]{ric05}). These models have been shifted up by 0.14 dex in Li abundance to make the initial abundance of the models, $\\log ({\\rm Li}/{\\rm H})=2.58$, coincide with the primordial Li abundance predicted from fluctuations of the microwave background measured by the WMAP satellite (\\cite[Cyburt et al. 2008]{cyburt}). The models assuming pure atomic diffusion, and, among those including turbulent mixing, T6.0 and T6.09, are ruled out by our observations. All such models predict that in dwarf stars Li should be either more abundant or the same as in subgiant stars. The only model which predicts a Li pattern which is qualitatively similar to that observed, is the T6.25 model.  Models including atomic diffusion and tachocline mixing (\\cite[Piau 2008]{pia08}) do not seem to reproduce our observations, since they provide a constant Li abundance up to 5500\\,K. The sophisticated models which besides diffusion and rotation also take into account the effect of internal gravity waves (\\cite[Talon \\& Charbonnel 2004]{tac04}), seem to predict accurately the Li abundance pattern in solar-type stars, at solar metallicity (\\cite[Charbonnel \\& Talon 2005]{cat05}), but models at low metallicity are still needed.  The cosmological lithium discrepancy still needs to be solved. Given that none of the existing models of Li evolution in stellar atmospheres matches the observations, we hope our results will prompt further new theoretical investigations."
    },
    "0910/0910.5713_arXiv.txt": {
        "abstract": "{ We present an analysis of the Sommerfeld enhancement to dark matter annihilation in the presence of an excited state, where the interaction inducing the enhancement is purely off-diagonal, such as in models of exciting or inelastic dark matter. We derive a simple and accurate semi-analytic approximation for the $s$-wave enhancement, which is valid provided the mass splitting between the ground and excited states is not too large, and discuss the cutoff of the enhancement for large mass splittings. We reproduce previously derived results in the appropriate limits, and demonstrate excellent agreement with numerical calculations of the enhancement. We show that the presence of an excited state leads to generically larger values of the Sommerfeld enhancement, larger resonances, and shifting of the resonances to lower mediator masses. Furthermore, in the presence of a mass splitting the enhancement is no longer a monotonic function of velocity: the enhancement where the kinetic energy is close to that required to excite the higher state can be up to twice as large as the enhancement at zero velocity.} ",
        "introduction": "In particle physics, perturbation theory is commonly employed to calculate annihilation and scattering cross sections, with higher-order terms in the perturbative expansion being neglected. Provided the theory is not strongly coupled, this is generally a good approximation for relativistic particles, but at low velocities and in the presence of a long-range force (classically, when the potential energy due to the long-range force is comparable to the particles' kinetic energy), the perturbative approach breaks down. In the nonrelativistic limit, the question of how the long-range potential modifies the cross section for short-range interactions can be formulated as a scattering problem in quantum mechanics, with significant modifications to the cross sections occuring when the particle wavefunctions are no longer well approximated by plane waves (so the Born expansion is not well-behaved). The deformation of the wavefunctions due to a Coulomb potential was calculated by Sommerfeld in \\cite{sommerfeld}, yielding a $\\sim 1/v$ enhancement to the cross section for short-range interactions (where the long-range behavior due to the potential can be factorized from the relevant short-range behavior).  Recent measurements of the positron and electron cosmic ray spectra by the PAMELA, ATIC, Fermi and H.E.S.S experiments have observed a rise in the positron fraction starting at $E \\sim 10$ GeV and extending (at least) up to $E \\sim 100$ GeV \\cite{Adriani:2008zr}, as well as a broad excess in the total $e^+ + e^-$ spectrum extending from several hundred GeV to $\\mathcal{O}(\\mathrm{TeV})$ energies \\cite{aticlatest, Abdo:2009zk, Aharonian:2009ah}. Observations by WMAP \\cite{Bennett:2003bz,Hinshaw:2008kr} also suggest an excess in microwave emission from the inner Galaxy, termed the ``WMAP Haze'' \\cite{Finkbeiner:2003im}, which has been attributed to synchrotron radiation from a new population of high energy electrons (with energies of tens to hundreds of GeV) \\cite{Finkbeiner:2004us, Dobler:2007wv, Hooper:2007kb}.  There has been considerable interest in these observations as possible signatures of dark matter (DM) annihilation or decay. If DM annihilation is responsible for the excesses, the annihilation cross section must exceed that required for a WIMP initially in thermal equilibrium to freeze out with the correct relic density, by 1-4 orders of magnitude depending on the dark matter mass. Furthermore, to avoid antiproton bounds from the PAMELA experiment \\cite{Adriani:2008zq} and to generate sufficiently hard $e^+ e^-$ spectra, the dark matter should annihilate largely into leptons, pions, kaons and other light states (e.g. \\cite{Cirelli:2008pk, Meade:2009iu}).  If dark matter is self-interacting, either via exchange of Standard Model gauge bosons or due to some novel interaction, then the nonperturbative ``Sommerfeld enhancement'' \\emph{must} be taken into account when computing annihilation cross sections in the present-day Galactic halo. The importance of the Sommerfeld enhancement in the context of dark matter annihilation has been studied  by several authors \\cite{Hisano:2003ec, Hisano:2004ds, Cirelli:2007xd, MarchRussell:2008yu}. Due to the velocity dependence of the Sommerfeld effect, and the presence of large resonances due to near-threshold bound states at low velocity (see e.g. \\cite{MarchRussell:2008tu}), the annihilation cross section could potentially be greatly increased in the present-day Galactic halo ($\\beta \\sim 10^{-4}-10^{-3}$), without greatly perturbing the freezeout of annihilations at $\\beta \\sim 0.3$ (however, the Sommerfeld enhancement can modify the relic density at some level, particularly in the presence of resonances, see e.g. \\cite{Dent:2009bv, Zavala:2009mi}).  Furthermore, it was pointed out in \\cite{Finkbeiner:2007kk,Cholis:2008vb} that the presence of light ($\\lesssim$ GeV) force carriers coupling to the dark matter would naturally lead to unusual dark matter annihilation signatures, with zero or small branching ratios into heavy hadrons and gauge bosons, and hard spectra of leptons, pions and other light states. The dominant annihilation channel in such models would be annihilation of the dark matter into the new force carriers, which would then decay into kinematically accessible SM states. In particular, provided the mass of the new force carrier was less than twice the proton mass, no excess antiprotons would be produced. A new GeV-scale force in the dark sector could therefore naturally generate both a large annihilation cross section via Sommerfeld enhancement, and the observed $e^+ e^-$ spectra, without violating antiproton constraints \\cite{ArkaniHamed:2008qn, Pospelov:2008jd}.  As noted in \\cite{ArkaniHamed:2008qn}, if dark matter is charged under some new dark gauge group, then when the new gauge boson acquires an $\\mathcal{O}$(GeV) mass the states in the dark matter multiplet can naturally be split from each other by a small ($\\mathcal{O}$(MeV)) amount (this happens automatically, due to loops of the new dark gauge bosons, if the gauge group is non-Abelian \\cite{ArkaniHamed:2008qn, Baumgart:2009tn}, but such mass splittings can also be naturally generated in Abelian models, e.g. \\cite{Cheung:2009qd}). Furthermore, if the dark matter is a Majorana fermion or real scalar, its couplings with any vector boson must be off-diagonal, since it cannot carry conserved charge. Models of this type, where the dark matter can only scatter inelastically at tree level (that is, by exciting some slightly more massive state) are also motivated by the iDM \\cite{Smith:2001hy, Chang:2008gd} and XDM \\cite{Finkbeiner:2007kk} scenarios, which explain, respectively, the anomalies observed by the DAMA/LIBRA \\cite{Bernabei:2008yi} and INTEGRAL/SPI \\cite{Weidenspointner:2006nu} experiments.  Despite the fairly generic presence of these excited states, most recent analyses of the Sommerfeld enhancement in the context of the cosmic-ray anomalies have treated only the case of degenerate dark matter states, under the presumption that the introduction of an excited state does not significantly modify the enhancement to annihilation, unless the mass splitting is large enough to cause the enhancement to cut off (see \\cite{ArkaniHamed:2008qn} for an argument to this effect) \\footnote{The related problem of collisional excitation of these states, via scattering mediated by a long-range force, has been considered by several authors, e.g. \\cite{Finkbeiner:2007kk,Chen:2009dm}.}. While it is true that the general qualitative features of the enhancement are mostly unchanged unless the mass splitting is quite large, modifications to the resonance structure occur for smaller mass splittings, and should be taken into account in any detailed analysis of the enhancement. However, the addition of even a single mass splitting adds an extra parameter to the problem, thus making numerical studies significantly more time-consuming; the numerical calculation of the enhancement can also become badly unstable in some parts of parameter space.  In \\S \\ref{sec:deriv} we derive simple expressions to aid in computing the Sommerfeld enhancement in the general case where the $l$th partial wave dominates annihilation, for an arbitrary potential coupling $N$ states. In \\S \\ref{sec:parametrics} we discuss the general properties of the case where the dark matter has a single excited state and the interaction is purely off-diagonal, including the limit of zero mass splitting and a characterization of the parts of parameter space where the enhancement is negligible. In \\S \\ref{sec:approxsol} we derive an approximate semi-analytic solution for the Schr\\\"{o}dinger equation in the case of a two-state system with purely off-diagonal interaction, focusing on the case of $s$-wave annihilation. In \\S \\ref{sec:analysis} we present a simple analytic approximation for the Sommerfeld enhancement, strictly valid when the annihilation channels and matrix elements are the same for two particles in the excited state as in the ground state, and also applicable to the more general case (up to a simple rescaling), provided the enhancement is large. We employ this approximation to discuss the new features present in the case with a non-zero mass splitting, and confirm our results by numerical simulations. ",
        "conclusions": "We have developed a detailed approximate expression for the $s$-wave Sommerfeld enhancement in the presence of a mass splitting, in the case where only inelastic scattering between the interacting particles is possible at tree level. The $\\delta \\rightarrow 0$ limit of our results yields an accurate approximation for the enhancement due to an attractive scalar Yukawa potential. The enhancement cuts off at low velocity if the mass splitting $\\delta \\gtrsim \\alpha^2 m_\\chi$, the effective Bohr energy of the two-body state, but even for smaller mass splittings, interesting modifications of the enhancement relative to the $\\delta \\rightarrow 0$ case can occur. Additionally, we have derived simple expressions to facilitate computation of the Sommerfeld enhancement due to an arbitrary $N \\times N$ matrix potential, where the $l$th partial wave dominates.  In the special case where all elements of the annihilation matrix are identical, our expression for the $s$-wave enhancement takes a simple form, \\[ S = \\frac{2 \\pi }{\\epsilon_v}\\sinh\\left(\\frac{\\epsilon_v \\pi }{\\mu}\\right) \\left\\{ \\begin{array}{cc} \\frac{1}{\\cosh\\left(\\epsilon_v \\pi/\\mu\\right)-\\cos\\left(\\sqrt{\\epsilon_\\delta^2-\\epsilon_v^2} \\pi /\\mu+2 \\theta_-\\right)} & \\quad \\epsilon_v < \\epsilon_\\delta, \\\\ \\\\ \\frac{\\cosh\\left(\\left(\\epsilon_v+\\sqrt{-\\epsilon_\\delta^2+\\epsilon_v^2}\\right) \\pi/2 \\mu\\right) \\text{sech}\\left(\\left(\\epsilon_v-\\sqrt{-\\epsilon_\\delta^2+\\epsilon_v^2}\\right) \\pi/2 \\mu \\right)}{ \\cosh\\left(\\left(\\epsilon_v+\\sqrt{-\\epsilon_\\delta^2+\\epsilon_v^2}\\right) \\pi /\\mu\\right)-\\cos(2 \\theta_-)} &\\quad \\epsilon_v > \\epsilon_\\delta. \\end{array} \\right. \\] where $\\mu$ and $\\theta_-$ are real functions which we have defined. This expression is valid provided $m_\\phi \\lesssim \\alpha m_\\chi$, $v/c \\lesssim \\alpha$, and $\\delta \\lesssim \\alpha^2 m_\\chi$, although it may become less accurate for $\\delta \\gtrsim \\max(\\alpha m_\\phi, \\, m_\\chi (v/c)^2)$. If any of the three former conditions are violated, there is no significant enhancement ($S \\sim \\mathcal{O}(1)$). In the regions where this expression is expected to be valid, comparison to numerical results demonstrates that it accurately reproduces all significant features of the enhancement. Furthermore, while derived for the case where all elements of the annihilation matrix are identical, a simple rescaling of this result by the average of the elements of the annihilation matrix (normalized to the $\\Gamma_{11}$ element) can accurately describe the enhancement for a general annihilation matrix.  Employing this approximate form for the enhancement, we have shown that below the threshold for on-shell production of the excited state, in comparison to the case with no mass splitting: \\begin{itemize} \\item The non-resonant, unsaturated enhancement is higher by a factor of $\\sim 2$. \\item The positions of the resonances shift to lower values of $\\epsilon_\\phi$ ($m_\\phi / \\alpha m_\\chi$) and become more widely spaced. The heights of the resonances increase by a factor of $\\sim 4$. \\item The enhancement $S$ is no longer a monotonic function of velocity: close to the threshold for on-shell excitation of the higher-mass state, the enhancement may be as much as a factor of $\\sim 2$ greater than its saturated value (i.e. its value in the limit $v \\rightarrow 0$). \\end{itemize}  In the context of DM annihilation, if the kinetic energy of WIMPs in the neighborhood of the Earth is close to the threshold for excitation of the higher-mass state (which is true by construction for models which realize the iDM mechanism), increased near-threshold enhancement may slightly alleviate constraints on DM annihilation from the early universe, where the DM is extremely slow-moving. The generically larger values of the unsaturated non-resonant enhancement, and the shift in position of the resonances, may help provide sufficiently large enhancements to generate recently observed cosmic-ray excesses, particularly for lower-mass force carriers. \\vskip 18pt \\noindent {\\bf Acknowledgements}: I am grateful to Nima Arkani-Hamed for suggesting this analysis, to Douglas Finkbeiner for many helpful discussions, and to Neal Weiner for pointing out useful references and the possibility of near-threshold resonances. I thank the anonymous referee for their useful feedback and suggestions; Lisa Goodenough, Rob Morris, David Poland, David Rosengarten, David Shih, Natalia Toro and Kathryn Zurek for helpful comments and conversations; and the NYU Center for Cosmology and Particle Physics and the Michigan Center for Theoretical Physics for their hospitality during the completion of this work. This work was partially supported by a Sir Keith Murdoch Fellowship from the American Australian Association.  \\appendix"
    },
    "0910/0910.0246_arXiv.txt": {
        "abstract": "This paper presents the results of collisional evolution calculations for the Kuiper belt starting from an initial size distribution similar to that produced by accretion simulations of that region - a steep power-law large object size distribution that breaks to a shallower slope at $r\\sim1-2$ km, with collisional equilibrium achieved for objects $r\\lesssim 0.5$ km. We find that the break from the steep large object power-law causes a divot, or depletion of objects at $r\\sim10-20$ km, which in-turn greatly reduces the disruption rate of objects with $r\\gtrsim 25-50$ km, preserving the steep power-law behavior for objects at this size. Our calculations demonstrate that the roll-over observed in the Kuiper belt size distribution is naturally explained as an edge of a divot in the size distribution; the radius at which the size distribution transitions away from the power-law, and the shape of the divot from our simulations are consistent with the size of the observed roll-over, and size distribution for smaller bodies. Both the kink radius and the radius of the divot center depend on the strength scaling law in the gravity regime for Kuiper belt objects. These simulations suggest that the sky density of $r\\sim1$ km objects is $\\sim10^6-10^7$ objects per square degree. A detection of the divot in the size distribution would provide a measure of the strength of large Kuiper belt objects, and constrain the shape of the size distribution at the end of accretion in the Kuiper belt. ",
        "introduction": "The Kuiper belt is a population of planetesimals outside the orbit of Neptune which exhibits high inclinations and eccentricities by many belt members \\citep{Trujillo2001b,Brown2001}, and has a very low mass, $M \\lesssim 0.1\\mbox{ M$_\\oplus$}$ \\citep{Gladman2001,Bernstein2004,Fuentes2008}. The high relative encounter velocities, $v_{rel} \\sim 1 \\mbox{ km s$^{-1}$}$ \\citep{Delloro2001} and infrequent collisions of the largest members \\citep{Davis1997,Durda2000} make the growth of Eris-sized bodies impossible over the age of the Solar system. Accretion in the early stages of planet-building must have been in a more dense and quiescent environment allowing large objects to grow \\citep{Stern1997a}.  The size distribution of the Kuiper belt is the result of the accretionary and collisional processes that have gone on in that region, and therefore provides one of the main constraints on those processes. For a general review of the Kuiper belt size distribution, see \\citet{Petit2008} \\footnote{Much of the observations that constrain the size distribution for radii $r\\sim20-60$ km are published in more recent works \\citep{Fraser2009,Fuentes2009}.} There are three properties which describe the general shape of the Kuiper belt size distribution: \\begin{itemize} \\item the existence of the largest objects, Eris and Pluto \\item the large object size distribution for $D\\gtrsim100$ km which is well characterized by a steep power-law $\\frac{dN}{dR}\\propto r^{-q}$ with $q\\sim4.8$ \\citep{Gladman2001,Petit2006,Fraser2008} \\item the existence of a ``roll-over'' at $r\\sim25-60$ km where the size distribution flattens to a much shallower distribution than for larger objects \\citep{Bernstein2004,Fuentes2008,Fraser2009} \\end{itemize} \\noindent These three properties need to be reproduced by any successful attempt at recreating the accretionary and collisional history of the Kuiper belt region. While this has not yet been achieved, these properties provide some insight into the general history of the belt.  Accretion simulations such as those of \\citet{Stern1996a,Kenyon2001} and \\citet{Kenyon2002} have demonstrated that objects as large as Eris could have accreted on timescales shorter than the age of the Solar system if the mass of the Kuiper belt was at least two orders of magnitude larger than the current mass, and if the relative encounter velocities, and hence, the eccentricities and inclinations of the objects were significantly lower in the early Solar system (see review by \\citet{Kenyon2008}). The existence of the steep large object size distribution however, suggests that accretion was a short-lived process \\citep{Kenyon2002,Fraser2008}. Some event must have disrupted accretion - likely the same event which scattered the primordial Kuiper belt onto orbits with high inclinations and eccentricities as observed today (see a review of proposed processes by \\citet{Morbidelli2008}) - otherwise the large object size distribution would be too shallow to be compatible with the observed distribution.  The simulations of \\citet{Stern1996a,Kenyon2001} and \\citet{Kenyon2002} cannot reproduce the large roll-over size. In these simulations, interaction velocities remain low, such that objects larger than $r\\sim 1$ km do not experience disruptive collisions over the age of the Solar system. Thus, the size distribution for objects larger than $r\\sim 2$ km exhibits a steep slope comparable to that observed today, and the accretion break, or roll-over, occurs at radii too small to be compatible with observations. This suggests that after the belt was dynamically excited, it remained massive enough for a time long enough to allow collisional erosion to produce the large roll-over observed today.  A model presented by \\citet{Kenyon2004} could reproduce the large roll-over if gravitational stirring by Neptune at its current location was included early in the simulations. Accretion occurred at a sufficient rate in these simulations to produce Eris-sized objects. In these simulations, accretion lasted too long however, resulting in a large object size distribution too shallow to be compatible with observations. In addition, the mass loss due to collisional grinding was insufficient to produce the tenuous modern-day belt (see discussions by \\citet{Morbidelli2008} and \\citet{Kenyon2008}). The Kuiper belt must have undergone a more rapid excitation and mass depletion than occurred in these calculations.  \\citet{Pan2005} presented an analytical, order of magnitude, collisional evolution model, reminiscent of the ground breaking work of \\citet{Dohnanyi1969}. With this model, they demonstrated that the large roll-over size could be produced by collisional grinding on timescales shorter than the age of the Solar system. They however, assumed that the size distribution rolled over to collisional equilibrium. It has been demonstrated that there is a range of objects with sizes smaller than the roll-over,  which are preferentially eroded but which are not being replenished from fragments of larger disrupted objects \\citep{Kenyon2002}. These objects do not achieve collisional equilibrium. Equilibrium is only achieved at a some smaller radius, the exact size of which, depends on the density of planetesimals. Thus, the model of \\citet{Pan2005} likely, does not predict the correct shape of the size distribution smaller than the roll-over, or the rate at which the roll-over evolves to larger sizes.  \\citet{Benavidez2009} present a collisional evolution model of the Kuiper belt, in which they include collisions from multiple dynamical classes. Using this model, they calculate the collisional evolution of the Kuiper belt, starting from an initially steep size distribution for all sizes, with slope similar to that observed for large objects. Their results confirm the findings of \\citet{Pan2005}; collisional grinding with relative velocities comparable to that observed can disrupt the majority of objects with $r\\sim50-100$ km, and in effect, produce a roll-over at that size. These calculations however, likely do not reproduce the collisional history of the Kuiper belt, as they start from an initial condition not expected from standard accretion processes \\citep{Kenyon2002}. Accretion simulations which produce Eris-sized objects produce size distributions with an accretion break at $r\\sim 1-2$ km, not a steep distribution for all sizes.   Here we present a collisional evolution model, which we utilize to calculate the size distribution of the Kuiper belt from some starting condition. With this model, we wish to determine whether or not collisional erosion could produce the $r\\sim20-50$ km roll-over, starting from a size distribution similar to that produced by models of accretion in the outer Solar system, but in the dynamically excited conditions of the current day Kuiper belt. From these calculations, we wish to constrain the shape of the modern day size distribution smaller than the roll-over, and to constrain the mass of the Kuiper belt at the end of accretion.   In Section~\\ref{sec:model} we present our collisional evolution model, and the planetesimal strength and shattering models we adopt. In Section~\\ref{sec:results} we present the main results of our calculation, namely that that existence of a break at $r\\sim1$ km left-over from accretion causes a divot in the size distribution at larger sizes. In Section~\\ref{sec:discussion} we present the consequences of our model, and finish with concluding remarks in Section~\\ref{sec:conclusions}. ",
        "conclusions": "} Our calculations have provided the first direct link from the ``typical'' size distribution produced from accretion, to the collisionally modified size distribution of the modern day Kuiper belt. Our calculations have demonstrated that, given a size distribution similar to that produced by accretion, the likely result of collisional evolution in the Kuiper belt, is to produce a divot in the Kuiper belt size distribution at $r\\sim10-15$ km, and a roll-over from the large object accretion slope at $r\\sim25-60$ km, compatible with the observed roll-over in the Kuiper belt size distribution. We conclude that the size distribution is likely not a power-law for objects smaller than the roll-over, and exhibits a dearth of objects at the divot location. Assuming a divot has been formed, then a measurement of the divot center radius would provide constraint on the strength scaling law in the gravity regime for KBOs, and constrain the size distribution of $r\\sim 1$ km objects left-over from the accretion processes in the early Solar system.   From our calculations, we find that the Kuiper belt must have had at least $1-5 \\mbox{ M$_\\oplus$}$ of material after its objects were excited onto orbits with eccentricities comparable to those observed, otherwise the roll-over would not be produced on Gyr timescales. If the belt had a mass of $\\sim 45 \\mbox{ M$_\\oplus$}$, then a divot would form in $\\gtrsim 40$ Myr.   From the results of our calculations, we predict a deep absence of objects at $r\\sim10$ km. We also predict  that the sky density of objects with $r\\sim1$ km is $\\sim10^6-10^7$ per square degree, 2-3 orders of magnitude larger than if the size distribution is a shallow power-law for objects smaller than the roll-over, as suggested by \\citet{Fraser2009}."
    },
    "0910/0910.4086.txt": {
        "abstract": "A new classification system for carbon-rich stars is presented based on an analysis of 51 AGB carbon stars through the most relevant classifying indices available \\citep{Keenan93,Morgan04}. The extension incorporated, that also represents the major advantage of this new system, is the combination of the usual optical indices that describe the photospheres of the objects, with new infrared ones, which allow an interpretation of the circumstellar environment of the carbon-rich stars. This new system is presented with the usual spectral subclasses and \\CC -, j-, MS- and temperature indices,  and also with the new SiC- (SiC/C.A. abundance estimation) and $\\tau$- (opacity) indices. The values for the infrared indices were carried out through a Monte Carlo simulation of the radiative transfer in the circumstellar envelopes of the stars. The full set of indices, when applied to our sample, resulted in a more efficient system of classification, since an examination in a wide spectral range allows us to obtain a complete scenario for carbon stars. ",
        "introduction": "The asymptotic giant branch stars with a ratio C$/$O $> 1$ have their optical spectra ruled by bands of carbon compounds, which obscure many atomic features. The green-red optical spectrum is dominated by Swan bands $^{12}$C$^{12}$C, Red System bands $^{12}$C$^{14}$N and sometimes present isotopic bands, e.g. $^{13}$C$^{12}$C and $^{13}$C$^{14}$N. As a result, a classification carried out through the optical atomic data on carbon stars is troublesome. During their ascent on the AGB phase, the mass loss from the star creates a circumstellar envelope of gas and dust. The compounds of this shell have their maximum emission on the infrared spectral region. The main infrared feature for carbon stars is the 11.3$\\mu$m emission due to the presence of SiC grains, which also represent an evidence of a C-rich dust envelope. Spitzer IRS spectra also confirmed two absorptions at 7 and 13.7 $\\mu$m, whose origin is still not well understood \\citep{speck06}.  The first classification system for carbon-rich stars was the Henry Draper Catalogue \\citep{HD}, in which the stars were presented in two spectral classes: type R and N, that were also divided into temperature subclasses. These subclasses were based on lines in the blue spectral region, that, although being a good source of information to determine the temperature classes of G and K stars, were not an appropriate choice for the cold carbon stars. Therefore, the distinction between an R8, the redder R-type stars, from an Na, the most blue N-type, was not obvious. \\citet{MK41} decided to rearrange all the carbon-rich stars in another classification scheme that presented a single temperature sequence. The numerical temperature type was, at this time, determined based on the atomic and molecular structures, considered more susceptible to temperature variations. They also added a \\CC -index, based on the strength of the \\CC~bands. Many attempts were made to improve this MK system, while keeping its basic structure. \\citet{Yama72,Yama75} with the C-Classification System listed in his tables a very large number of stars with the two main parameters from the MK System, and additional intensity indices of several other atomic and molecular characteristics from carbon stars. The result of this attempt, was a detailed notation for carbon-rich objects, but it was not very practical or compact.  Later, \\citet{Keenan93} pointed out several reasons for replacing this old C-Classification System, since much importance was given to the Na D-line in determining the numerical indices of temperature. The assumption that the Na D-line could be a good tracer of temperature in this case, however, was revealed to be quite unfortunate, specially in the case of N-type stars, as they have an enormous molecular opacity in the same spectral region due to the (7,2) CN band. Furthermore, the N-type and the R-type stars, in fact, describe two different populations that should not be classified under the same spectral class. The new classification system proposed by Keenan93, the MK~Revised System, re-established the spectral subclasses for the carbon rich objects and the temperature indices based on infrared intensities. Additionally, this Revised MK System listed four abundance indices from \\citet{Yama72}: the intensity of the \\CC~band, the isotopic carbon ratio, the Si\\CC\\ band and the CH band strength. These indices and system are the most widely accepted and used nowadays in the study of carbon stars \\citep{Morgan03,Morgan04,deroo07}. %%tentar achar uma referencia atual????  Concerning the circumstellar analysis, \\citet{sloan98} presented the first classification based on an infrared study of the spectra of carbon and oxygen-rich AGB stars. Qualitative classes for the circumstellar material were settled through two indices: one chemical and other a strength indicator of the characteristic structure of the emission spectrum.  In the present paper we present a new scheme of classification for C-rich AGB stars (hereafter NSCC) that includes a notation conceived through a wide analysis of the main spectroscopic features of these carbon-rich objects with circumstellar material in its surroundings. The main difference between this scheme and the others available is that, instead of using a single region of the spectrum, we suggest a classification system based on a large wavelength range, from the blue-optical to the mid-infrared. Thus, this notation draws a more complete scenario of each star. The parameters used in this extended classification system were devised from the analysis of a sample of 51 observed AGB carbon stars.  Section \\ref{obs} describes the optical observations and other infrared data used and section \\ref{optical} and \\ref{ir} details how indices and parameters of classification were obtained. While, section \\ref{optical} explains the optical indices, section \\ref{ir} contains the infrared ones. Section \\ref{disc} shows an analysis of the classification indices, a discussion of possible evolutionary sequences for the AGB carbon stars and also an atlas of optical spectra and infrared radiative transfer models of the stars studied in this work. The overall results are discussed in section \\ref{conc}. ",
        "conclusions": "\\label{conc}  This New Scheme of Classification of C-Rich AGB Stars can not be applied to any carbon star. It is not a good application for AGB carbon stars that have a very thick envelope, e.g. extreme AGB stars can not be treated this way as their optical spectra are highly obscured. Moreover, the methodology employed is not tied to the sample presented, it has the flexibility to serve to all these kind of carbon stars by using the given coefficients and parameters. The indices correspond to either a direct measurement of the intensities and equivalent widths of features observed in low resolution spectra of carbon-rich stars or obtained through a radiative transfer model fit to infrared data.  The seven indices presented describe in detail the complex scenario of the carbon rich stars. It is possible just by analyzing the compact final notation to get a full set of basic information about an AGB carbon star. As all calibrations were established based on the well quoted works, the indices and levels employed represent the more successful historical parameters on the study of carbon stars.  Regarding the use of the indices, we demonstrated how the SiC/A.C. ratio and the $\\tau$ parameters could together provide information about the evolutionary stages of the carbon-rich stars. The evolutionary sequences here proposed may provide a new insight over the nature of objects such as the odd C-J stars.  %% The \\notetoeditor{TEXT} command allows the author to communicate %% information to the copy editor.  This information will appear as a %% footnote on the printed copy for the manuscript style file.  Nothing will %% appear on the printed copy if the preprint or %% preprint2 style files are used.   %###############################################"
    },
    "0910/0910.2243_arXiv.txt": {
        "abstract": "We have developed a homogeneous model of physical chemistry to investigate the neutral-dominated, water-based Enceladus torus.  Electrons are treated as the summation of two isotropic Maxwellian distributions$-$a thermal component and a hot component.  The effects of electron impact, electron recombination, charge exchange, and photochemistry are included.  The mass source is neutral H$_2$O, and a rigidly-corotating magnetosphere introduces energy via pickup of freshly-ionized neutrals.  A small fraction of energy is also input by Coulomb collisions with a small population ($<$\\,1\\%) of supra-thermal electrons.  Mass and energy are lost due to radial diffusion, escaping fast neutrals produced by charge exchange and recombination, and a small amount of radiative cooling.  We explore a constrained parameter space spanned by water source rate, ion radial diffusion, hot-electron temperature, and hot-electron density.  The key findings are: (1) radial transport must take longer than 12 days; (2) water is input at a rate of 100--180 kg s$^{-1}$; (3) hot electrons have energies between 100 and 250 eV; (4) neutrals dominate ions by a ratio of 40:1 and continue to dominate even when thermal electrons have temperatures as high as $\\approx$\\,5 eV; (5) hot electrons do not exceed 1\\% of the total electron population within the torus; (6) if hot electrons alone drive the observed longitudinal variation in thermal electron density, then they also drive a significant variation in ion composition. ",
        "introduction": "Absorption of UV starlight during occultation of Saturn's moon Enceladus showed that it continuously ejects neutral H$_2$O at a rate of $\\approx$\\,150--300\\,kg\\,s$^{-1}$ from water-ice geysers located at its southern pole \\citep{2006Sci...311.1422H}.  Models suggest that the water and its chemical by-products form an extended neutral-dominated torus centered on the orbit of Enceladus \\citep{jurac2005}.  Similarly, Jupiter's volcanic moon Io emits a mixture of SO$_2$ and S$_2$ at a rate of $\\approx$\\,1\\,ton\\,s$^{-1}$, and chemical by-products produce a plasma torus centered on the orbit of Io (see review by \\cite{thomas04}).  The Hubble Space Telescope (HST) observations by \\cite{1993Natur.363..329S} revealed that the neutral-to-ion ratio in the Enceladus torus ($\\approx$\\,10) is three orders of magnitude greater than in the Io torus ($\\approx$\\,10$^{-2}$).  Compositional differences and the degree of ionization within these two systems can be attributed to their chemistry (Io's based on sulfur dioxide and Enceladus's based on water) as well as the fact that fresh ions are picked with five times more energy in the Io torus than in the Enceladus torus \\citep{2007GeoRL..3409105D}.  An important lesson learned from studying the physical chemistry of the Io torus is that a small fraction of hot electrons ($<$\\,1\\%) play a critical role in determining composition.  To model the Cassini UltraViolet Imaging Spectrograph (UVIS) data obtained during Cassini's E2 flyby of Jupiter (October 2000 to March 2001), \\cite{steffl2,steffl1,steffl3,steffl4} adapted the \\cite{2003JGRA..108.1276D} Io torus model.  \\textit{Steffl et al.}\\ used their models to study radial, temporal, and azimuthal variation in mixing ratios (ion-to-electron density ratios), thermal-electron density, and thermal-electron temperature.  \\cite{steffl4} concluded that hot electrons are necessary for the Io torus energy budget and that two modulations of the hot-electron population are required to reproduce both the temporal and spatial variations in composition observed in the data, one modulating in Jupiter's System III longitude, the other in System IV.  We anticipate that hot electrons are similarly important in Saturn's Enceladus torus.  \\cite{2007GeoRL..3409105D} developed a simplified oxygen-based model to compare the Enceladus and Io tori.  They found that collisional heating by a population of hot electrons is much less important at Enceladus, contributing only 0.5\\% of the energy to the torus, compared to 60\\% at Io.  They also cited two major reasons for the discrepancy in the neutral-to-ion ratio between the two systems.  First, newly created ions are picked up in the Io torus by Jupiter's magnetosphere at roughly five times the energies as are those in the Enceladus torus by Saturn's magnetosphere.  The higher-energy pickup ions in the Io torus warm the thermal electrons, which then reduces the neutral-to-ion ratio via impact ionization.  Second, because of the high abundance of molecular ions compared with atomic oxygen ions in the Enceladus torus (e.g., \\cite{sittler2005}), \\cite{2007GeoRL..3409105D} expected that molecular dissociative recombination (not included in their model) will therefore drive the ratio even higher in the Enceladus torus.  From the conclusion of \\cite{2007GeoRL..3409105D}:  ``The addition of the full water-group molecular chemistry will introduce an additional plasma sink through dissociative recombination of the molecular ions.  Therefore, our simplified O-based chemistry likely represents a lower limit for the neutral/ion ratio.'' To investigate the consequences of a water-based Enceladus torus dominated by molecular chemistry, we have improved on the Delamere model by including a comprehensive set of water-based reactions and species to more accurately estimate steady-state densities and temperatures of ions and electrons.  We also add neutral and ionized molecules.  Molecules are more abundant in the Enceladus torus where low plasma density allows H$_2$O to escape from Enceladus largely intact, whereas the more energetic local plasma interaction at Io results in dissociation of SO$_2$ \\citep{2008JGRA..11309208D}.  Moreover, thermal electrons ($\\approx$\\,2\\,eV) throughout the Enceladus torus dissociate H$_2$O approximately ten times less easily than thermal electrons ($\\approx$\\,5\\,eV) in the Io torus dissociate SO$_2$ (\\textit{V.\\ Dols, personal comm.}).  Previous models of molecular chemistry in Saturn's magnetosphere were driven by Pioneer 11 and Voyager 1 and 2 observations \\citep{frank1980,trainor1980,wolfe1980,bridge1981,bridge1982,sittler1983}.  \\cite{richardson1986} showed the importance of recombination under conditions of slow radial transport.  Using essentially the same chemical reactions as \\cite{richardson1986}, \\cite{1998JGR...10320245R} determined ion and neutral lifetimes within Saturn's inner magnetosphere ($\\lesssim 12$$R_\\mathrm{S}$; $R_\\mathrm{S}\\equiv$ Saturn radius $ =6.0\\times10^9\\,\\mathrm{cm}$), constrained by HST observations of the extended OH cloud (see their Figure 2, and references therein).  They solved the rate equations for number densities while also solving the radial diffusion equation, but energy conservation was not considered.  \\cite{2002GeoRL..29x..25J} and \\cite{jurac2005} further improved on these models by considering neutral cloud expansion, and solved for plasma and neutral distributions self-consistently in order to study the source of water within Saturn's inner magnetosphere.  Our model concentrates on the molecular chemistry.  We start with a uniform box and characterize transport by just a time scale.  However, we do consider energy balance.  More importantly, unlike the above models, we retain H$_3$O$^+$ in our model, which proves to be a significant component.  The purpose of this paper is to summarize the sensitivity of the chemistry of Enceladus's torus to several parameters.  The parameters we  investigated are hot-electron temperature, hot-electron density, H$_2$O source rate, ion radial-diffusion time scale, and proton temperature anisotropy.  Being lighter, protons are less bound to the equator \\citep{1980GeoRL...7...41B,2008P&SS...56....3S,2009JGRA..11404211P}.  This means protons spend only a fraction of the time interacting with the heavy ions and molecules.  This effect is simulated with a `proton dilution' factor (Section \\ref{modl}).  We search for values of these parameters leading to thermal-electron temperature, thermal-electron density, and water-group ($\\mathrm{W^ +\\equiv O^++OH^++H_2O^++H_3O^+}$) ion-to-proton ratio consistent with available Cassini data (Section \\ref{constraints}).  The observations used to constrain our model and to define the parameter space are given in Section \\ref{data}.  The model is described in Section \\ref{modl}.  The best fit (baseline solution) and the procedure used to find it are discussed in Section \\ref{BaselineSolution}.  Model sensitivity to each of the parameters is discussed in Section \\ref{grdSrch}.  Finally, the importance of hot electrons with regard to water-group ion composition is demonstrated in Section \\ref{fehModulation}. ",
        "conclusions": "We have compared output from our model to the constraints on thermal-electron density ($n_\\mathrm{e}$), thermal-electron temperature ($T_\\mathrm{e}$), and mixing ratio of water-group ions to protons (W$^+/$H$^+$) by exploring the space spanned by the following four parameters:  neutral source rate ($N_\\mathrm{src}$), hot-electron temperature ($T_\\mathrm{eh}$), hot-electron density ($n_\\mathrm{eh}\\equiv f_\\mathrm{eh}n_\\mathrm{e}$), and radial transport time scale ($\\tau_\\mathrm{trans}$).  Our important results are:  \\begin{enumerate} \\item For the constraint choices of $n_\\mathrm{e}=60\\ \\mathrm{cm}^{-3}$, $T_\\mathrm{e}=2$ eV, and W$^+$/H$^+$=12, we find the following limits on the parameters: \\begin{enumerate} \\item[\\ ] $1.5\\lesssim N_\\mathrm{src}/(10^{-4}\\mathrm{\\ cm^{-3}\\ s^{-1}})\\lesssim 2.7$ \\item[\\ ] $12\\ \\mathrm{days} \\lesssim \\tau_\\mathrm{trans}$ \\item[\\ ] $0.3\\lesssim f_\\mathrm{eh}(\\%)\\lesssim0.9$ \\item[\\ ] $100\\lesssim T_\\mathrm{eh}/\\mathrm{eV}\\lesssim250$. \\end{enumerate} \\noindent  For a volume of $2\\pi(4R_\\mathrm{S})(2R_\\mathrm{S})^2$, the source rate can be scaled to give a mass source rate of  $100\\lesssim N_\\mathrm{src}/(\\mathrm{kg\\ s^{-1}})\\lesssim 180$.  We find that $f_\\mathrm{H^+}$ (the fraction of protons confined to the equator) is strongly coupled to the hot-electron population and has not been constrained by this study (Section \\ref{varyingProtonDilution}). The solution space can be compared with future measurements of the parameters ($N_\\mathrm{src},\\ f_\\mathrm{eh},\\ T_\\mathrm{eh},\\ \\tau_\\mathrm{trans}$) and composition mixing ratios.  Upper limits on UV power emanating from the Enceladus torus from neutral oxygen at 1304, 1356, and 6300\\,\\AA\\  would be very useful.  \\item With the full water-based chemistry, photo- plus impact ionization is nearly as important as charge exchange at providing energy by way of fresh pickup ions (Figure \\ref{enerFlow}). \\item The water-based chemistry (particularly dissociative recombination) increases the neutral-to-ion ratio from 12 (the \\cite{2007GeoRL..3409105D} oxygen-based model) to $\\approx 40$.  Further, the Enceladus torus remains neutral-dominated even when the thermal-electron temperature approaches the temperature of electrons in Jupiter's Io torus (6 eV). \\item The H$_3$O$^+$/W$^+$ ratio is directly correlated with the W$^+$/H$^+$ ratio (Figure \\ref{barchart}), implying that H$_3$O$^+$ is strongly anti-correlated with H$^+$ (Section \\ref{alternativeConstraints}).  This result is important because obtaining the W$^+$/H$^+$ ratio is more straightforward than obtaining individual water-group abundances from the CAPS data.  However, \\cite{2008P&SS...56....3S} have shown, based on CAPS SOI data, that H$_3$O$^+$ dominates the water-group near the orbit of Enceladus (their Figure 15).  The dominance of H$_3$O$^+$ seen in the CAPS data has not been obtained by our model with the given constraints.  It is likely (manuscript in preparation) that the local interaction of the corotating plasma with the Enceladus plumes may contribute significantly to the  H$_3$O$^+$ abundance \\citep{cravens2009}. \\item Hot electrons are necessary in the Enceladus to enhance ionization torus but do not directly contribute more that 1\\% of the total electron population. \\item Significant variation in composition can be driven by a small perturbation in the hot-electron population (Figure \\ref{fehVoutput}). \\end{enumerate} The sensitivity study presented here will be useful in a future interpretation of longitudinal and temporal observations (e.g., \\cite{gurnett2007}) of the Enceladus torus, especially in the context of a modulating hot-electron density."
    },
    "0910/0910.5180_arXiv.txt": {
        "abstract": "Galaxies, in particular disc galaxies, contain a number of structures and substructures with well defined morphological, photometric and kinematic properties. Considerable theoretical effort has been put into explaining their formation and evolution, both analytically and with numerical simulations.  In some theories, structures form during the natural evolution of the galaxy, i.e. they are a result of nature. For others, it is the interaction with other galaxies, or with the intergalactic medium -- i.e. nurture -- that accounts for a structure. Either way, the existence and properties of these structures reveal important information on the underlying potential of the galaxy, i.e. on the amount and distribution of matter -- including the dark matter -- in it, and on the evolutionary history of the galaxy. Here, I will briefly review the various formation scenarios and the respective role of nature and nurture in the formation, evolution and properties of the main structures and substructures. ",
        "introduction": "What do we mean by ``nature versus nurture''? To some degree, all galaxies have one or more neighbours and/or are influenced by their surroundings. So in the absolute sense of the word, there is no such thing as an isolated galaxy, except in computer simulations. On the other hand, even though a given structure may be triggered by an interaction, it will evolve within its galaxy potential, which, together with the remaining characteristic properties of the galaxy, will influence its evolution. Thus, no structure is 100\\% due to nature, or 100\\% due to nurture. It is, nevertheless, useful to ask which of the two has mainly influenced the formation and evolution of a given structure, and, therefore, in the following I will assign the origin of a structure to one of these two alternatives.  In fact, as will be seen below, more than one scenario is possible for each structure, and, very often, some mixture of nature and nurture can be involved. Hence, one can only reason in terms of probabilities and propose this or that agent as the most probable cause for a given feature. Time comes into play as well: if one goes sufficiently far back, all galaxies have experienced important interactions with their surroundings. If, however, the time elapsed since the last interaction is longer than the time necessary for a given structure to form, one can safely argue that this structure is due to nature rather than to nurture (see also a discussion in \\cite{HopkinsKMQ09b}).  After these introductory remarks, let me review what models and simulations can tell us. Since the subject is very broad, I will discuss here only certain aspects, necessarily reflecting my own preferences, without me wanting to belittle in any way work which I cannot mention for lack of time and space. The list of references given is likewise far from complete. Finally, I will only discuss the formation of structures and not their destruction, although the latter could also be due either to nature or to nurture. ",
        "conclusions": "I discussed the possible formation mechanisms of structures and substructures in galaxies, focusing on whether their origin is due to nature or to nurture. In most cases, there is more than one alternative. A few structures can be explained only by a mechanism relying on nurture, but none uniquely only to nature."
    },
    "0910/0910.1595_arXiv.txt": {
        "abstract": "We have undertaken a deep ($\\sigma\\sim 1.1$ mJy) 1.1-mm survey of the $z=0.54$ cluster MS\\,0451.6$-$0305 using the AzTEC camera on the James Clerk Maxwell Telescope. We detect 36 sources with S/N$\\ge 3.5$ in the central 0.10 deg$^2$ and present the AzTEC map, catalogue and number counts. We identify counterparts to 18 sources (50\\%) using radio, mid-infrared, \\Spitzer\\ IRAC and Submillimeter Array data. Optical, near- and mid-infrared spectral energy distributions are compiled for the 14 of these galaxies with detectable counterparts, which are expected to contain all likely cluster members. We then use photometric redshifts and colour selection to separate background galaxies from potential cluster members and test the reliability of this technique using archival observations of submillimetre galaxies. We find two potential MS\\,0451$-$03 members, which, if they are both cluster galaxies have a total star-formation rate (SFR) of $\\sim100$ \\myr\\ -- a significant fraction of the combined SFR of all the other galaxies in MS\\,0451$-$03. We also examine the stacked rest-frame mid-infrared, millimetre and radio emission of cluster members below our AzTEC detection limit and find that the SFRs of mid-IR selected galaxies in the cluster and redshift-matched field populations are comparable. In contrast, the average SFR of the morphologically classified late-type cluster population is $\\sim 3$ times less than the corresponding redshift-matched field galaxies. This suggests that these galaxies may be in the process of being transformed on the red-sequence by the cluster environment. Our survey demonstrates that although the environment of \\ms\\ appears to suppress star-formation in late-type galaxies, it can support active, dust-obscured mid-IR galaxies and potentially millimetre-detected LIRGs. ",
        "introduction": "\\label{sec:intro}   Galaxy clusters are highly biased environments in which galaxies potentially evolve more rapidly than in the field. The galaxy populations of local massive clusters contain mainly early-type galaxies which define a colour-magnitude relation (CMR) \\citep{Visvanathan77, Bower92}. However, studies of clusters out to $z\\sim 1$ suggest that they contain increasing activity at higher redshifts due to a growing fraction of blue, star-forming galaxies \\citep{Butcher84}. Over the same redshift range there appears to be a growing deficit in the CMR population at faint magnitudes, as well as a claimed increasing decline in the numbers of S0 galaxies, suggesting that the blue, star-forming galaxies may be transforming into these passive populations with time \\citep{Dressler97, Smail98a, DeLucia07, Stott07, Holden09}.  The blue, star-forming populations within the clusters are accreted from the surrounding field as the clusters grow via gravitational collapse. The evolution in the star-forming fraction in the clusters may thus simply track the increasing activity in the field population at higher redshifts. The increasing activity in the field is also reflected in an increasing number of the most luminous (and dusty) star-burst galaxies with redshift \\citep[e.g.][]{LeFloch05}: the Luminous Infrared Galaxies (with L$_{\\rm FIR}\\geq 10^{11}$\\,L$_\\odot$) and their Ultraluminous cousins (ULIRGs, L$_{\\rm FIR}\\geq 10^{12}$\\,L$_\\odot$). These systems will also be accreted into the cluster environment along with their less-obscured and less-active population as the clusters grow.  Indeed, mid-infrared (mid-IR) surveys of clusters have identified a population of dusty star-bursts whose activity increases with redshift  \\citep[e.g.][]{Geach06, Geach08}. However, these mid-IR surveys can miss the most obscured (and potentially most active) galaxies which are optically thick in the rest-frame mid-IR. If they are present in clusters -- even in relatively low numbers -- such active galaxies will contribute significantly to the star formation rate (SFR) in these environments and the metal enrichment of the intracluster medium. Hence to obtain a complete census of the star formation within clusters we need to survey these systems at even longer wavelengths, in the rest-frame far-infrared (far-IR), corresponding to the observed sub-millimetre and millimetre waveband.  Over the past decade or more there have been a number of studies of rich clusters of galaxies in the sub-millimetre and millimetre wavebands \\citep[e.g.][]{Smail97, Zemcov07, Knudsen08}. Many of these studies were seeking to exploit the clusters as gravitational telescopes to aid in the study of the distant sub-millimetre galaxy (SMG) population behind the clusters, while others focused on the detection of the Sunyaev-Zel'dovich (SZ) emission. Due to the limitations of current technologies direct detection of millimetre sources is restricted to those with fluxes S$_{\\rm 1100\\umu m}\\ga 1$\\,mJy, or equivalently galaxies with SFRs $\\ga 300$\\,M$_{\\sun}$\\,yr$^{-1}$ -- much higher than expected for the general cluster populations based on optical surveys. Nevertheless these studies have serendipitously detected a number of cluster galaxies, although these are either atypical (e.g.\\ central cluster galaxies, \\citet{Edge99}) or are not confirmed members \\citep[e.g.][]{Best02, Webb05}. More critically,  with few exceptions these studies have all focused on the central 2--3 arcmin of the clusters, where the SZ emission and lensing amplification both peak, and have not surveyed the wider environment of the cluster outskirts where much of the obscured star formation is likely to be occurring \\citep[e.g.][]{Geach06}. The two exceptions are the wide-field survey of the $z\\sim0.25$ cluster A\\,2125 by \\citet{Wagg09} and the survey of an overdense region of the COSMOS field by \\citet{Scott08} and \\citet{Austermann09a}. \\citet{Wagg09} detected 29 millimetre galaxies across a $\\sim 25$-arcmin diameter region centered on A\\,2125 of which none are claimed to be members. However, the only redshift estimates available are based on crude radio-to-millimetre spectral indices, which are sensitive to both dust temperature and redshift \\citep{Blain03}. The AzTEC/COSMOS survey \\citep{Scott08, Austermann09a} covered a number of structures, including an X-ray detected $z=0.73$ cluster and concluded that the statistical overdensity of sources was most likely due to the gravitational lensing of background SMGs by these foreground structures.  To help to conclusively answer the issue of the obscured star-forming populations in distant clusters we have therefore undertaken a wide-field 1.1-mm survey of the $z=0.54$ rich cluster MS\\,0451.6$-$0305 (hereafter \\ms) with the AzTEC camera  \\citep{Wilson08} on the James Clerk Maxwell Telescope (JCMT). This panoramic millimetre survey can also take advantage of the significant archival data available for this well-studied region. In particular the panoramic imaging of \\ms\\ from {\\it Spitzer} and uniquely, the {\\it Hubble Space Telescope} ({\\it HST}), as well as extensive archival multiwavelength imaging and spectroscopy, which we employ for determining cluster membership of AzTEC-detected galaxies. Our survey probes the rest-frame 700\\,$\\umu$m emission of cluster galaxies -- in search of examples of strongly star-forming but obscured galaxies -- as well as identifying more luminous and more distant examples of the SMG population.   We can compare our millimetre search for cluster members to the previous mid-IR survey of this cluster by \\citet{Geach06} who uncovered a population of dusty star-forming galaxies which dominate the integrated SFR of the cluster of $200\\pm 100$\\,M$_{\\sun}$\\,yr$^{-1}$. Our AzTEC map covers $\\sim 60$ times the area of the 850\\,$\\umu$m SCUBA observations of the central region of \\ms\\ \\citep{Borys04a}, while the depth of $\\sigma\\sim1.1$ mJy is sufficient to identify ULIRGs individually and obtain stacked constraints on far-IR fainter cluster members.  We describe our observations and the data reduction in \\S\\ref{sec:obsred}; present 1.1-mm number counts, identify counterparts to the AzTEC galaxies, and use photometric techniques to determine cluster membership in \\S\\ref{sec:analysis}. \\S\\ref{sec:conc} contains our conclusions; individual AzTEC galaxies are presented in Appendix \\ref{sourcenotes}. Throughout this paper we use J2000 coordinates and $\\Lambda$CDM cosmology with $\\Omega_M=0.3$, $\\Omega_{\\Lambda}=0.7$ and $H_0=70$ ${\\rm kms^{-1}Mpc^{-1}}$ and all photometry is on the AB system unless otherwise stated. ",
        "conclusions": "\\label{sec:conc}  In this study we have investigated the dust-obscured star-forming population of the galaxy cluster \\ms; we utilise 1.1-mm observations to study obscured star-formation in \\ms. We present a $\\sigma\\sim 1.1$ mJy AzTEC map of the central 0.1 deg$^2$ of \\ms, within which 36 sources are detected at S/N$\\ge3.5$. We use radio, 24\\,$\\umu$m, IRAC and SMA observations to precisely locate 18 of these SMGs.  We calculate the reliability of photometric redshifts for SMGs and find that they are able to remove the bulk of background contamination, allowing us to isolate potential cluster members using our optical, near- and mid-IR photometry. Based on these redshifts we find two SMGs which are possible cluster members: MMJ\\,045421.55 and MMJ\\,045431.35. These systems are both resolved into close pairs of galaxies by our ground-based and \\HST\\ imaging, suggesting that interactions have triggered their starbursts. If they are cluster members both of these SMGs contain cold dust with $T_d\\sim30$ K, for $\\beta=1.1$ \\citep[similar to SLUGS galaxies;][]{Dunne00a} and have SFRs of $\\sim50$ \\myr\\ each.  \\citet{Geach06} compared obscured activity, based on $24 \\umu$m emission, in \\ms\\ and \\cl\\ at $z=0.39$, and found that \\cl\\ has ${\\rm SFR}\\sim5$ times that of \\ms\\ within 2 Mpc of the clusters centres. They show that, taking into account the slightly different redshifts and masses of these two structures, \\ms\\ is underactive and \\cl\\ overactive at $24 \\umu$m. Therefore, it is likely the other galaxy clusters, including \\cl, could contain ULIRGs in significantly larger numbers than we find in \\ms.  To further investigate the obscured star-forming population which lies below the limit of our AzTEC observations we create composite SEDs of spectroscopically confirmed mid-IR, early- and late-type cluster members and compare them to the corresponding field populations. As expected we find that both early-type populations are undetected at both 24\\,$\\umu$m and 1.1 mm and so are unlikely to be actively forming large numbers of stars ($<0.1$ \\myr). The  24\\,$\\umu$m galaxies are significantly more active than the morphologically classified late-types with SFRs $\\sim15$ \\myr\\ versus $\\sim3$ \\myr\\ on average. We find that the star-formation activity in the cluster late-type population, compared to a redshift-matched field population is quenched and $\\sim3$ times lower than expected. Mid-IR galaxies do not show this trend suggesting the more intense activity in these systems is more robust to environmental influences. We find that the total ${\\rm SFR>315\\pm50}$ \\myr\\ in the central 2 Mpc of \\ms. However, if MMJ\\,045421.55 is a cluster member it has ${\\rm SFR}\\sim50$ \\myr\\ and lies 1 Mpc from the cluster centre, taking the total SFR within 2 Mpc to $\\ga360$ \\myr."
    },
    "0910/0910.0962_arXiv.txt": {
        "abstract": "In the context of coronal heating, among the zoo of MHD waves that exist in the solar atmosphere, Alfv\\'en waves receive special attention. Indeed, these waves constitute an attractive heating agent due to their ability to carry over the many different layers of the solar atmosphere sufficient energy to heat and maintain a corona. However, due to their incompressible nature these waves need a mechanism such as mode conversion (leading to shock heating), phase mixing, resonant absorption or turbulent cascade in order to heat the plasma. Furthermore, their incompressibility makes their detection in the solar atmosphere very difficult. New observations with polarimetric, spectroscopic and imaging instruments such as those on board of the japanese satellite \\textit{Hinode}, or the \\textit{SST} or \\textit{CoMP}, are bringing strong evidence for the existence of energetic Alfv\\'en waves in the solar corona. In order to assess the role of Alfv\\'en waves in coronal heating, in this work we model a magnetic flux tube being subject to Alfv\\'en wave heating through the mode conversion mechanism. Using a 1.5-dimensional MHD code we carry out a parameter survey varying the magnetic flux tube geometry (length and expansion), the photospheric magnetic field, the photospheric velocity amplitudes and the nature of the waves (monochromatic or white noise spectrum). The regimes under which Alfv\\'en wave heating produces hot and stable coronae is found to be rather narrow. Independently of the photospheric wave amplitude and magnetic field a corona can be produced and maintained only for long ($> 80$ Mm) and thick (area ratio between photosphere and corona $> 500$) loops. Above a critical value of the photospheric velocity amplitude (generally a few km s$^{-1}$) the corona can no longer be maintained over extended periods of time and collapses due to the large momentum of the waves. These results establish several constraints on Alfv\\'en wave heating as a coronal heating mechanism, especially for active region loops. ",
        "introduction": "New observations from polarimetric, spectroscopic and imaging instruments are revealing a corona permeated with waves. Compressive modes such as the slow and fast magnetohydrodynamic (MHD) modes have a long chapter in observational history \\citep[see reviews by ][]{Nakariakov_2005LRSP....2....3N, Banerjee_2007SoPh..246....3B, Ruderman_2009SSRv..tmp...54R, Taroyan_2009SSRv..tmp...24T}. Only recently, the development of high resolution instruments has brought strong evidence for the existence of the third MHD mode, the Alfv\\'en mode, in the solar atmosphere. In the majority of the observational reports of Alfv\\'en waves ambiguity exists and the waves can be interpreted as trapped modes, fast and slow kink waves. \\citet{DePontieu_2007Sci...318.1574D} reported transversal displacements of spicules from observations in the Ca II H-line (3968 \\AA) with a broadband filter of \\textit{Hinode}/SOT. The wavelengths of these chromospheric agents were estimated to be longer than 20000 km, with periods between 100 and 500 s and speeds of at least 50-200 km s$^{-1}$. In the absence of a stable waveguide they interpreted the swaying of the spicules as a result of upward propagating Alfv\\'en waves. However, it was pointed out by  \\citet{Erdelyi_2007Sci...318.1572E} that these waves were likely to be kink oscillations, due to the displacement of the axis of symmetry of the flux tube caused by the later. Using the Coronal Multi-Channel Polarimeter (\\textit{CoMP}) \\citet{Tomczyk_2007Sci...317.1192T} analyzed properties of infrared coronal emission line Fe\\,XIII (1074.7 nm) across a large field of view with short integration times. Waves propagating along magnetic field lines with ubiquitous quasi-periodic fluctuations in velocity with periods between 200 and 400 s (power peak at 5-min) and negligible intensity variations were reported. Wavelengths and phase speeds were estimated to be higher to 250 Mm and 1 Mm s$^{-1}$ respectively. In this work it was suggested that these waves were Alfv\\'en waves probably being generated in the chromospheric network from mode conversion of p-modes propagating from the photosphere (hence explaining the peak in the power spectrum of the waves). On the other hand \\citet{Erdelyi_2007Sci...318.1572E} and \\citet{VanDoorsselaere_2008ApJ...676L..73V} argued that an interpretation in terms of fast kink waves was more appropriate due mainly to the observed collective behavior which should be absent in the case of Alfv\\'en waves. Being compressible waves it was argued that kink waves would appear however incompressible in the corona for an instrument such as \\textit{CoMP} due to the very small variation in intensity they produce.  The difficulty in detecting the Alfv\\'en wave is due in part to its incompressible nature. Being only a transverse disturbance in the magnetic field, it is practically invisible to imaging instruments, unless the structure of propagation is displaced periodically from the line of sight \\citep{Williams_2004ESASP.547..513W}. Also, due to the large wavelengths (a few megameters) and short periods (a few minutes) that may be involved in the corona, a large field of view, short integration time and proper resolution are needed for their detection in the corona. With polarimeters and spectrographs however Alfv\\'en waves are more easily detected, as suggested by \\citet{Erdelyi_2007Sci...318.1572E}. These waves propagate as torsional disturbances of the magnetic field, causing a periodic and spatially dependent spectral line broadening \\citep{Zaqarashvili_2003AA...399L..15Z}. Up to date, the only observational report on torsional Alfv\\'en waves is the recent work by \\citet{Jess_2009Sci...323.1582J}, in which observations in H$\\alpha$ of a bright point using the Solar Optical Universal Polarimeter (SOUP) of the \\textit{SST} were reported. No variation in intensity nor line-of-sight velocity was detected and a periodic spectral broadening of the H$\\alpha$ line was found, leaving Alfv\\'en waves as the only possible interpretation for their results. The waves were reported to have periods in the 100-500 s interval (with maximum power in the 400-500 s range). By comparing H$\\alpha$ continuum with H$\\alpha$ core images they reported a canopy-like structure with magnetic fields expanding $\\sim1300$ km in a vertical distance of $\\sim1000$ km. This large expansion is crucial for Alfv\\'en wave heating theories. For instance, assuming an Alv\\'en speed of 10 km s$^{-1}$ in the upper photosphere/ lower chromosphere and a period of 100 s (a value in the lower range of periods reported in the works cited above) we obtain a wavelength of 1000 km for Alfv\\'en waves propagating from the photosphere, matching the distance of large expansion of the flux tubes. Moreover, the height in the atmosphere where this takes place is the region of transition from high beta to low beta plasma. We thus have ideal conditions for mode conversion. Trapped modes will thus exchange their energy and a wide range of waves with different characteristics are likely to issue from this region.  The large emphasis that has been put on the search for Alfv\\'en waves in the solar corona is largely due to their connection with coronal heating. Theoretically, they can be easily generated in the photosphere by the constant turbulent convective motions, which inputs large amounts of energy into the waves \\citep{Muller_1994AA...283..232M, Choudhuri_1993SoPh..143...49C}. Having magnetic tension as its restoring force the Alfv\\'en wave is less affected by the large transition region gradients with respect to other modes. Also, when traveling along thin magnetic flux tubes they are cut-off free since they are not coupled to gravity (Musielak et al. 2007)\\footnote{\\citet{Verth_etal_aap_09} have pointed out however that the assertion made by \\citet{Musielak_2007ApJ...659..650M} is valid only when the temperature in the flux tube does not differ from that of the external plasma. When this is not the case a cut-off frequency is introduced.}. Alfv\\'en waves generated in the photosphere are thus able to carry sufficient energy into the corona to compensate the losses due to radiation and conduction, and, if given a suitable dissipation mechanism, heat the plasma to the high million degree coronal temperatures \\citep{Uchida_1974SoPh...35..451U, Wentzel_1974SoPh...39..129W, Hollweg_1982SoPh...75...35H, Poedts_1989SoPh..123...83P, Ruderman_1997AA...320..305R, Kudoh_1999ApJ...514..493K} and power the solar wind \\citep{Suzuki_2006JGRA..11106101S, Cranmer_2007ApJS..171..520C}.  Another possible generation mechanism for Alfv\\'en waves is through magnetic reconnection. The amount of energy imparted to these waves during the reconnection process may depend on the location in the atmosphere of the event and is a subject of controversy. \\citet{Parker_1991ApJ...372..719P} suggested a model in which 20~\\% of the energy released by reconnection events in the solar corona is transfered as a form of Alfv\\'en wave. \\citet{Yokoyama_1998ESASP.421..215Y} studied the problem simulating reconnection in the corona, and found that less than 10~\\% of the total released energy goes into Alfv\\'en waves. This result is similar to the 2-D simulation results of photospheric reconnection by \\citet{Takeuchi_2001ApJ...546L..73T}, in which it is shown that the energy flux carried by the slow magnetoacoustic waves is one order of magnitude higher that the energy flux carried by Alfv\\'en waves. On the other hand, recent simulations by Kigure et al. (private communication) show that the fraction of Alfv\\'en wave energy flux in the total released magnetic energy during reconnection in low $\\beta$ plasmas may be significant (more than 50~\\%). Since the observed ubiquitous intensity bursts (nanoflares) are thought to play an important role in the heating of the corona \\citep{Hudson_1991SoPh..133..357H} and since they are generally assumed to be a signature of magnetic reconnection it is then crucial to determine the energy going into the Alfv\\'en waves during the reconnection process. Moreover, it is possible that the observed nanoflares themselves are a consequence of Alfv\\'en wave heating as porposed by\\citet{Moriyasu_2004ApJ...601L.107M}. In that work the resulting intensity bursts producing the nanoflares are created by Alfv\\'en waves by first converting to longitudinal modes which steepen into shocks and heat the plasma. This model was further developed by \\citet{Antolin_2008ApJ...688..669A} (hereafter, Paper 1) and \\citet{Antolin_2009arXiv0903.1766A}, where it was shown that the frequency of the resulting heating events from Alfv\\'en waves followed a power law distribution. It was further shown that Alfv\\'en wave heating could be differentiated from nanoflare-reconnection heating during observations through a series of signatures \\citep[see also][]{Taroyan_2008IAUS..247..184T, Taroyan_2009SSRv..tmp...24T}. For instance, Alfv\\'en waves lead to a dynamic, uniformly heated corona with steep power law indexes (issuing from statistics of heating events) while nanoflare-reconnection heating leads to lower dynamics (unless catastrophic cooling takes place in the case of footpoint concentrated heating) and shallow power laws.  The main problem faced by Alfv\\'en wave heating is to find a suitable dissipation mechanism. Being an incompressible wave it must rely on a mechanism by which to convert the magnetic energy into heat. Several dissipation mechanisms have been proposed, such as parametric decay \\citep{Goldstein_1978ApJ...219..700G, Terasawa_1986JGR....91.4171T}, mode conversion \\citep{Hollweg_1982SoPh...75...35H, Kudoh_1999ApJ...514..493K, Moriyasu_2004ApJ...601L.107M}, phase mixing \\citep{Heyvaerts_1983AA...117..220H, Ofman_2002ApJ...576L.153O}, or resonant absorption \\citep{Ionson_1978ApJ...226..650I, Hollweg_1984ApJ...277..392H, Poedts_1989SoPh..123...83P, Erdelyi_1995AA...294..575E}. The main difficulty lies in dissipating sufficient amounts of energy in the correct time and space scales. For more discussion regarding this issue the reader can consult for instance \\citet{Klimchuk_2006SoPh..234...41K, Erdelyi_2007AN....328..726E} and \\citet{Aschwanden_2004psci.book.....A}.  Due to the observed increasing importance of Alfv\\'en waves in the solar atmosphere in this work we address the subject of coronal heating by Alfv\\'en waves in which mode conversion and parametric decay are taken as dissipation mechanisms. We concentrate on the heating of closed magnetic structures which populate the solar atmosphere, such as coronal loops, and analyze the efficiency of this heating model by carrying out a parametric space survey. We take a 1.5-dimensional model and we consider different photospheric drivers (a random driver creating a white noise spectrum, or a monochromatic driver in which several periods are tested), vary the photospheric magnetic field, the loop expansion from the photosphere to the corona and the loop length, and determine the regimes for which Alfv\\'en wave heating plays an important role in coronal heating. The paper is organized as follows. In \\S\\ref{two} we set up the loop geometries and define the equations. As in Antolin et al. (2008) the loops are modeled with a 1.5-dimensional MHD code including thermal conduction and radiative cooling. In \\S\\ref{three} we analyze the effect of varying the parameters on the thermodynamic structure of the loop. Further discussion and conclusions are presented in \\S\\ref{four}. ",
        "conclusions": "\\label{four}  Alfv\\'en waves constitute one of the main candidates for heating the solar corona. Theoretically easy to be generated from convective turbulent motions in the photosphere it has been shown that these waves are able to carry sufficient amounts of energy in order to balance energy losses from radiation and conduction and heat the corona up to the observed few million degrees \\citep{Uchida_1974SoPh...35..451U, Wentzel_1974SoPh...39..129W, Hollweg_1982SoPh...75...35H, Poedts_1989SoPh..123...83P, Ruderman_1997AA...320..305R, Kudoh_1999ApJ...514..493K}. Finding a suitable dissipation mechanism is a very difficult task, which has spawned a large active research community in solar physics. In this work we have considered a model in which mode conversion acts as the dissipative mechanism. More precisely, through nonlinear effects based on density fluctuations and wave-to-wave interaction in the chromosphere and corona the Alfv\\'en mode is able to transfer some part of its energy to the fast and slow modes, which steepen into shocks and heat the plasma.  We have conducted a parameter survey in which the effect of geometrical quantities such as the loop expansion and the loop's length on the efficiency of Alfv\\'en wave heating is studied. We further investigate the effect on the heating of other physical parameters such as the photospheric magnetic field and the generation of Alfv\\'en waves with a monochromatic and white noise spectrum in the photosphere.  In Fig.~\\ref{fig8} we plot the volumetric heating rate with respect to the considered parameters of the loop: period of the monochromatic longitudinal waves issuing from mode conversion of Alfv\\'en waves (top left panel), photospheric magnetic field (top right panel), expansion of the loop (bottom left panel), and loop length (bottom right panel). The volumetric heating rate for each case corresponds to its extremum values in the range between $\\langle v_{\\phi}^{2}\\rangle^{1/2}\\sim1.3$  and 2.3 km s$^{-1}$ found in the e panels of Figs.~\\ref{fig3}, \\ref{fig5}, \\ref{fig6} and \\ref{fig7}, which is the most efficient amplitude range for coronal heating we have found. We can see that the volumetric heating rate is an increasing function for all the considered parameters in the chosen azimuthal photospheric velocity interval. This figure summarizes well the results of our parameter survey. As the amplitudes of the twists in the photosphere generating the Alfv\\'en waves increase the momentum and energy flux of the waves increase. We have found that nonlinear effects generally increase as well, and mode conversion happens not only in the chromosphere but ubiquitously in the corona, thus heating the plasma. Suitable conditions for Alfv\\'en wave heating are thus strongly dependent on the nonlinear effects. The importance over nonlinearity favors thick and long loops (bottom left and right panels), and strong photospheric magnetic fields (top right panel). The dependence on the later is however not crucial and flux tubes with 1 kG field concentrations at their footpoints may be heated efficiently by Alfv\\'en waves. On the other hand, the expansion and the length of the loop are crucial parameters for the heating. As we can see in Fig.~\\ref{fig8}, the volumetric heating rate in the bottom left and right panels does not vary much respect to the loop expansion and the length, which shows the high sensitivity of the thermodynamics in the loop on the heating. Nonlinearity is directly proportional to the loop expansion. Furthermore, the strong expansion of loops from the photosphere to the transition region happens in a height comparable to the wavelength of the Alfv\\'en wave, which leads to the deformation of the Alfv\\'en wave and excitation of slow modes \\citep{Moore_1991ApJ...378..347M}. Since the plasma $\\beta$ parameter in that region is close to unity, further mode conversion is expected. Alfv\\'en waves have long wavelengths and in order for them to convert into longitudinal waves which steep into shocks they need to propagate relatively long distances. Loops with lengths lower than 80 Mm cannot be heated by Alfv\\'en waves in the present model. Caution must be taken however when drawing conclusions from these results. When calculating the lengths of loops in observations what is actually calculated is the length of the corona, which is the visible part of loops in EUV or X-ray images. In our model, the lengths of the coronae with average temperatures of $\\sim1$ MK we obtain with Alfv\\'en wave heating oscillate between 50 and 80 Mm with a mean around 70 Mm. Furthermore, we have defined a stable corona, as a corona which can be maintained over roughly 5.7 hours. However, observations show that coronal loops are dynamical entities which exhibit heating and cooling processes constantly. Likewise, the loops with a collapsing corona obtained here exhibit heating events with nanoflare-like temperatures constantly, which can match the observed bursty X-ray intensity profiles and the statistics displaying power law distribution of the heating events (cf. Paper 1). Also, a loop which is visible in soft X-rays over a large period of time may in fact result from the effect of different threads in the loop being heated at different times, thus giving the impression of a steady and uniformly heated loop \\citep[as discussed in][]{Patsourakos_2009ApJ...696..760P}. In the simple model we have considered we have assumed that the heating acts uniformly along the radial direction, thus disabling the possibility of many threads being heated at different times. We thus cannot completely reject the possibility of Alfv\\'en waves heating the coronae of short and thin loops. This idea should however be tested in an extended (at least 2.5 dimensional) version of this model.  All coronae show a range of velocities for which they are stable throughout the simulation and for which their lengths are constant. However, when the photospheric velocity field exceeds a critical value the corona collapses. The critical velocity is dependent on the parameters of the model. In this regime wave heating can no longer account for the large cool mass flux from the increasing wave amplitudes and collapses. Generally, the collapse is not due to a local increase in the corona of the radiative losses from the formation of a cool condensation as in the case of catastrophic cooling \\citep{Antiochos_1999ApJ...512..985A}, and does not exhibit ``limit cycles\" \\citep{Muller_2003AA...411..605M}. The later are characteristic of the catastrophic cooling mechanism which is proposed as an explanation for coronal rain. Further investigation in that direction is however required and will be the subject of a future paper.  All of the obtained coronae are uniformly heated irrespective of the parameters in the model. This is a characteristic of Alfv\\'en wave heating. The distribution of the heating in coronal loops as derived from observations is a matter of debate. Indeed, different methods for the analysis of observational data may lead to different results. A famous example of this fact is the analysis of the data set from soft X-ray observations with the \\textit{Yohkoh}/SXT instrument. An interpretation of the heating distribution of the observed loops in terms of apex concentrated heating \\citep{Reale_2002ApJ...580..566R}, footpoint concentrated heating \\citep{Aschwanden_2001ApJ...559L.171A} and uniform heating \\citep{Priest_1998Natur.393..545P} has been given. There seems to be however more observational evidence of coronal loops being heated towards the footpoints in active regions \\citep{Aschwanden_2001ApJ...560.1035A, Aschwanden_2000ApJ...541.1059A}. Further evidence of this fact has been found by \\citet{Hara_2008ApJ...678L..67H} using the \\textit{Hinode}/EIS instrument, in which active region loops are shown to exhibit upflow motions and enhanced nonthermal velocities in the hot lines of Fe XIV 274 and Fe XV 284. Possible unresolved high-speed upflows were also found. In Paper 1 we found that footpoint or uniform nanoflare heating exhibit hot upflows, thus fitting in the observational scenario of active regions, while Alfv\\'en wave heating was found to exhibit hot downflows, which may fit in the observational scenario of quiet Sun regions  \\citep{Chae_1998ApJS..114..151C, Brosius_2007ApJ...656L..41B}. The uniform heating from Alfv\\'en waves further supports this conclusion, since it is a characteristic of loops more generally found in quiet Sun regions. Active region loops exhibit low expansion factors due to the high magnetic field filling factor in those regions. In this chapter we have found that Alfv\\'en wave heating is effective only in thick loops (with area expansions between photosphere and corona higher than 600), further emphasizing our conclusion that active regions may not be heated by Alfv\\'en waves. These results seem to point to an important role of Alfv\\'en wave heating in Quiet Sun regions, where loops are often long, expand more than in active regions, and kG (or higher) bright points are ubiquitous.  Another important result we have found is that both long (above 100 s) and short period (below 50 s) waves (where the period is the resulting period of the longitudinal modes from mode conversion; the period of the Alfv\\'en waves is twice the stated period) play an important role in the heating of the loops. Long period waves produce very hot heating events in the corona and thus increase the average temperature of the corona due to the strong shocks they produce. Short period waves are responsible for keeping the corona, thus making it stable through uniform heating from the numerous weak shocks they produce. We have therefore a compromise between long and short period waves leading to efficient Alfv\\'en wave heating. The 300 s period waves are found to have considerable power to heat the corona even with low photospheric velocity amplitudes of $\\sim0.2$ km s$^{-1}$. For these waves resonances seem to be triggered in the loop, which allow higher transmission from the waves into the corona and subsequent mode trapping leading to efficient dissipation. This resonant damping mechanism is however not observed for random wave generation (white noise spectrum) or with other geometric parameters for the loop. A thorough study of these coronal resonances for Alfv\\'en waves will be the subject of a future paper."
    },
    "0910/0910.2226_arXiv.txt": {
        "abstract": "{Recent progress in instrumentation enables solar observations with high resolution simultaneously in the spatial, temporal, and spectral domains. } % {We use such high-resolution observations to study small-scale structures and dynamics in the chromosphere of the quiet Sun. } % {We analyse time series of spectral scans through the \\ion{Ca}{II}~854.2~nm spectral line obtained with the CRISP instrument at the Swedish 1-m Solar Telescope. The targets are quiet Sun regions inside coronal holes close to disc-centre. } % {The line core maps exhibit relatively few fibrils compared to what is normally observed in quiet Sun regions outside coronal holes. The time series show a chaotic and dynamic scene that includes spatially confined ``swirl'' events. These events feature dark and bright rotating patches, which can consist of arcs, spiral arms, rings or ring fragments. The width of the fragments typically appears to be of the order of only 0\\,\\farcs2, which is close to the effective spatial resolution. They exhibit Doppler shifts of $-2$ to $-4$\\,km$/$s but sometimes up to $-7$\\,km$/$s, indicating fast upflows. The diameter of a swirl is usually of the order of 2\\,\\arcsec. At the location of these swirls, the line wing and wide-band maps show close groups of photospheric bright points that move with respect to each other. }  % {A likely explanation is that the relative motion of the bright points twists the associated magnetic field in the chromosphere above. Plasma or propagating waves may then spiral upwards guided by the magnetic flux structure, thereby producing the observed intensity signature of Doppler-shifted ring fragments. } ",
        "introduction": "Advances in observational performance have revolutionised our understanding of the solar chromosphere over the last years. It is in particular the combination of high spectral, temporal, and spatial resolution that let us discover details and phenomena hitherto unaccessible. Our picture of the solar atmosphere changed from a static plane-parallel stratification to a very complex compound of intricately coupled dynamic domains (see recent reviews by e.g. Schrijver 2001; Judge 2006; Rutten 2006, 2007;Wedemeyer-B{\\\"o}hm et al. 2009). And yet many details concerning the chromosphere and its connection to the layers above and below remain elusive. The progress in our understanding is certainly hampered by the complexity and accessibility of currently available diagnostics probing that atmospheric layer. Among them are the spectral lines of \\ion{Ca}{II}, which are formed over a very extended height range in the atmosphere. Only the line cores truly originate from above the photosphere. Observing the line core exclusively therefore requires filters with a very narrow transmitted wavelength range, as the width of the line cores is only of the order of 100\\,pm or less.  Considering the spectral, temporal, and spatial domains, instrumental limitations enforce compromises which until recently only allowed high resolution in one or two of the domains, at the cost of the remaining one(s). This situation has now improved substantially with the installation of fast two-dimensional spectropolarimetric imagers at solar telescopes with large aperture. Examples are the IBIS instrument (Cavallini 2006) at the Dunn Solar Telescope, CRISP (Scharmer et al. 2008) at the Swedish Solar Telescope (SST), and the G{\\\"o}ttingen Fabry-P{\\'e}rot (Puschmann et al. 2006) Achieving high resolution simultaneously in all these three domains is like opening a new observing window on the chromosphere. Here we report on the discovery of small but fast rotating swirls in the chromosphere. There are quite a few examples known of rotating or vortex-like motions on the Sun: on large scales in the form of rotating sunspots (e.g., Brown et al. 2003), on smaller scales like vortices in the photospheric granulation (Brandt et al. 1988), and on the smallest scales in the form of whirlpool motion by magnetic bright points in the inter-granular lanes (Bonet et al. 2008). This rotating motion receives considerable interest since these photospheric motions have a profound effect on the outer-atmospheric magnetic fields as they have their roots in the photoshere. The stresses that are built up through rotation of magnetic fields are linked to the onset of solar flares and coronal heating (e.g., Parker 1983).  After this introduction we describe the observations in Sect.~\\ref{sec:observ}, which are analysed in Sect.~\\ref{sec:ana}. A discussion and conclusions are given in Sect.~\\ref{sec:discus}. ",
        "conclusions": "\\label{sec:discus}  At first glance, the swirls in the coronal hole are reminiscent of the convectively driven vortex flows reported by Bonet et al. (2008). In a high-resolution SST G-band filtergram time sequence, they found examples of groups of BPs displaying clear ``whirlpool'' motion. Bonet et al. associate the photospheric whirlpools with vigorous downdrafts at the vertices of intergranular lanes that have been predicted by convection models. The G-band BPs act as flow tracers which follow spiral trajectories when they are engulfed by the downdrafts. The size of these whirlpools is less than an arcsecond. Contrastingly, the photospheric BPs associated with the chromospheric swirls in our data do not follow clear spiral trajectories. It is therefore unclear if photospheric whirlpools and chromospheric swirls are connected.  Our wide-band images usually show a small number of photospheric BPs that move in relation to each other, buffeted by the lateral granular flows at the top of the convection zone. Being confined to the intergranular lanes, the BPs can eventually come close and even move past each other very closely. Our interpretation is that these close encounters result in a deformation and twist of the magnetic field continuing above the BPs. While tending to unwind the build-up stresses, gas may spiral upwards, producing the observed intensity signature of Doppler-shifted ring fragments.  Such a process may be impeded by the presence of a stronger magnetic field component in the chromosphere in the form of a ``magnetic canopy'' (e.g., Schrijver \\& Title 2003) as it exists above quiet Sun regions outside coronal holes. Swirls are therefore expected to be significantly less frequent in quiet Sun regions outside coronal holes, which have a high fibril coverage. Quantifying this assumption is unfortunately difficult as the detection of swirls in Ca line core maps is complicated by the presence of fibrils. Events ongoing in the layer just below the fibrils could be obscured by the horizontal component of the magnetic field (Vecchio et al. 2009).  The exact cause for the swirl phenomenon remains to be found. The magnetic fields, which are supposedly responsible for forcing the rotating plasma or propagating waves on ring-like or spiral trajectories, are not directly detected in our CRISP observations. However, all our detected swirls show a small group of BPs in the centre, implying a causal connection.  We interpret these chromospheric swirl motions and associated BP motions as a direct indication of upper-atmospheric magnetic field twisting and braiding as a result of convective buffeting of magnetic footpoints. This mechanism is one of the prime candidates for coronal heating (Parker 1988)."
    },
    "0910/0910.0015_arXiv.txt": {
        "abstract": "In a recent paper we presented the first semi-analytic model of galaxy formation in which the Thermally-Pulsing Asymptotic Giant Branch phase of stellar evolution has been fully implemented. Here we address the comparison with observations, and show how the TP-AGB recipe affects the performance of the model in reproducing the colours and near-IR luminosities of high-redshift galaxies. We find that the semi-analytic model with the TP-AGB better matches the colour-magnitude and colour-colour relations at $z \\sim 2$, both for nearly-passive and for star-forming galaxies. The model with TP-AGB produces star-forming galaxies with red V-K colours, thus revising the unique interpretation of high-redshift red objects as 'red \\& dead'. We also show that without the TP-AGB the semi-analytic model fails at reproducing the observed colours, a situation that cannot be corrected by dust reddening. We also explore the effect of nebular emission on the predicted colour-magnitude relation of star-forming galaxies, to conclude that it does not play a significant role in reddening their colours, at least in the range of star-formation rates covered by the model. Finally, the rest-frame K-band luminosity function at $z \\sim 2.5$ is more luminous by almost 1 magnitude. This indicates that the AGN feedback recipe that is adopted to regulate the high-mass end of the luminosity function should be sophisticated to take the effect of the stellar populations into account at high redshifts. ",
        "introduction": "Hierarchical clustering is the favoured scenario to describe the formation and evolution of matter structures in the universe (White \\& Rees, 1978), and semi-analytic models of galaxy formation proved themselves to be a powerful tool of investigation since the first formulation (White \\& Frenk, 1991). Over the years, many such models have been developed (see for instance Balland et al. 2003, Baugh et al. 2005, Bower et al. 2006, Cattaneo et al. 2008, Cole et al. 2000, De Lucia et al. 2004, Hatton et al. 2003, Kauffmann et al. 1993, Menci et al. 2006, Monaco et al. 2007, Somerville et al. 2008). The successes and failures of these models are strictly linked to those of the hierarchical scenario itself, ultimately depending on the mechanisms of mass accretion of objects as a function of time. The large-scale structure and the integrated properties of the galaxy population (such as the total stellar mass density for instance) are well reproduced. The detailed evolution of galaxies however presents several puzzling aspects, such as e.g. the size of the disks of spirals (Burkert 2009), or the $\\alpha$-enhancement and the ages of the stellar populations in massive ellipticals (Thomas 1999, Thomas et al. 1999, Cimatti et al. 2004, Nagashima et al. 2005, Thomas et al. 2005, Pipino et al. 2009, Kormendy et al. 2009). Moreover, global properties of the galaxy population such as the evolution of the stellar mass function (Cole et al. 2001, Bell et al. 2003, Bundy et al 2006-2009, Pozzetti et al. 2009, Colless et al. 2007) are still not reproduced in the models (Bundy et al. 2007, Marchesini et al. 2009, Kajisawa et al. 2009, Kodama \\& Bower 2003), although there is controversy on this point (Drory et al. 2004, Benson et al. 2007). It is debated if any of these problems can possibly be resolved with enhanced resolution in the simulations and more sophisticated recipes in the models.  One of the most problematic issues for the models is to reproduce the abundance of high redshift luminous galaxies (e.g., Conselice et al. 2007, Cimatti et al 2004, van Dokkum et al. 2004, 2006). This difficulty is partly due to a mis-interpretation of the nature of these objects.  The high-luminosity end of the galaxy population up to redshift $z \\sim 2.5$ consists in fact of objects that look like the early-type galaxies in the local universe, \\textit{i.e.} they are characterized by very red colours in the optical and near-IR, and high near-IR luminosities (Mancini et al. 2009, Cimatti et al 2004, McCarthy et al. 2004, Daddi et al. 2005, Saracco et al. 2005, Kriek et al. 2006). Local ellipticals showing the same photometry are old (with stellar populations older than $\\sim1$ Gyr), passively evolving, and with stellar masses $M_{star} > 10^{11} \\ M_{\\odot}$. Moreover, newer observations are building the case for the presence of extremely red and IR-luminous objects at even higher redshifts (Rodighiero et al. 2007, Mancini et al. 2009, Fontana et al. 2009).  The problem posed by the presence of these high redshift ($z>2$) red and luminous galaxies stems from the consensus that they are massive objects evolving passively, the so-called 'red \\& dead' galaxies. With the stellar population models currently used in the semi-analytic models in the literature (for the most part Bruzual \\& Charlot 2003, hereafter BC03), the only way to explain the high near-IR luminosities of high-redshift galaxies is to advocate very high galaxy masses and very old ages of the stellar populations. But these are not achieved in the actual model realizations (except at low redshifts), because the hierarchical mass assembly has an intrinsic difficulty in putting together massive and old objects at early epochs. In fact, the hierarchical scenario predicts a steady decline of the abundance of massive galaxies with increasing redshift (van Dokkum et al. 2004).  In Tonini et al. (2009) we showed that the predictions of colours and luminosities of galaxies at high redshift in a semi-analytic model are greatly affected by the recipes in use for the stellar populations, expecially the inclusion of the Thermally-Pulsing Asymptotic Giant Branch (TP-AGB). As shown in M05, in stellar populations of intermediate age ($\\leq 0.2 - 2$ Gyr) the TP-AGB phase dominates the near-IR luminosity, with a contribution up to $80 \\%$ in the rest-frame K band, and contributes to up to $40 \\%$ of the bolometric luminosity (M05). High-redshift galaxies, in which the mean age of the stellar populations is in that range, are expected to be dominated by the TP-AGB emission in the near-IR. This has been recently confirmed by SED-fitting of observations made with the Spitzer Space Telescope (Maraston et al. 2006, Cimatti et al. 2008). In Tonini et al. (2009) we included a complete treatment of the TP-AGB phase in the semi-analytic model GalICS (Hatton et al. 2003), by implementing the M05 stellar population models into the code, and showed that the rest-frame $V-K$ colours at high redshift get redder by more than 1 magnitude. Relatedly, the K-band mass-to-light ratio is shifted towards luminosities 1 magnitude higher for a given galaxy mass. Notably, actively evolving, star-forming high-redshift galaxies are predicted to have $V-K$ colours and near-IR luminosities similar to those of local, passively evolving massive systems (Tonini et al. 2009).  Once the stellar emission is correctly modeled with an exhaustive treatment of all the significant phases of stellar evolution, a more accurate comparison between the semi-analytic model predictions and the data is possible. In particular, the performance of the semi-analytic model in reproducing the observed colours and luminosities in the near-IR becomes meaningful to test the hierarchical mass assembly at different redshifts.  In the literature the comparison between galaxy formation models and data is typically done by obtaining physical properties for the real objects through application of stellar population models to data. However, this approach carries several degeneracies, including the adopted population synthesis model, the recipe for star formation history, the choice of metallicity, etc. When a realistic errorbar including all these variables ia attached to the observationally derived quantity, such as in Marchesini et al. 2009 (and see also Conroy et al. 2009 for a discussion), the results of such comparisons may not be clear-cut.  In this paper we adopt a different philosophy for the comparison between model and data. Instead of using processed data in the rest-frame system, we consider raw, unprocessed, apparent magnitudes straight out of the catalogues. We then produce mock catalogues out of the simulation, so that the output spectra of the model galaxies are redshifted in the observer's frame. The model apparent magnitudes and colours can then be directly compared with the observational data. This comparison yields direct information about the physical quantities in the model in use.  This procedure is straightforward and does not add substantial degeneracy that can jeopardize the comparison. A degeneracy that clearly remain is how dust reddening affects the intrinsic stellar emission, as recently pointed out by Guo \\& White (2008) and Conroy et al. (2009). However, we shall show that considering the intrinsic star formation rates in the model and using data mapping the rest-frame near-IR, such an uncertainty plays actually a minor role.  The structure of the paper is as follows. In Section 2 we briefly introduce the new semi-analytic model GalICS with the TP-AGB implementation through the M05 models (as from Tonini et al. 2009). In Section 3 we describe the data samples used for our analysis. In Section 4 we compare the colour-magnitude and colour-colour relations predicted by the model against samples of $z \\sim 2$ galaxies. In Section 4 we compare the model rest-frame K-band luminosity function in the M05 and Pegase cases with the predictions by other semi-analytic models. In Section 5 we discuss our results. ",
        "conclusions": "In a recent work (Tonini et al. 2009) we introduced the complete treatment of the TP-AGB phase of stellar evolution into a semi-analytic model of galaxy formation, by inserting the Maraston (2005) SSP models into the code GalICS. In the work presented here we compared the predictions on the near-IR luminosities and colours of high-redshift galaxies with data samples of nearly-passive and star-forming galaxies around $z \\sim 2$. Our main results are:  $\\cdot$ the TP-AGB is fundamental to allow the semi-analytic model to reproduce the observed optical and near-IR colours of both nearly-passive and star-forming galaxies at $z \\sim 2$; the inclusion of the TP-AGB increases the Irac3 luminosity (rest-frame K) and shifts the H-Irac3 (rest-frame V-K) colours by more than 1 magnitude;  $\\cdot$ without the TP-AGB, it is not possible to match the observed galaxy colours and luminosities by a modification of the dust reddening recipe alone;  $\\cdot$ the TP-AGB emission does not alter the optical luminosity and colours of star-forming galaxies. On the other hand, star formation does not dilute the TP-AGB emission in the near-IR. Even star-forming galaxies, very blue in the optical, can be very red in the near-IR. Therefore the labelling of red galaxies as 'red and dead' is misleading;  $\\cdot$ the nebular emission, produced by young stellar populations, does not add a significant contribution to the colours of star-forming galaxies, in the range of star-formation rates covered by the model; for SFR$>70 \\ M_{\\odot}/yr$ the rest-frame V-K colour is reddened by 0.2-0.4 mags;  $\\cdot$ the predicted mass-luminosity relation is affected by the inclusion of the TP-AGB; for a given galaxy mass, the rest-frame K-band luminosity is higher by more than 1 mag at $z>1$. As a consequence, the K-band luminosity function predicted by the model with the TP-AGB shifts redwards, expecially at the high-mass end, for $z>1$ (by $\\sim 0.7$ mag at $z \\sim 2.5$). The spread in the luminosity function between runs with and without the TP-AGB is comparable to the scatter caused by different AGN-feedback recipes in the literature.  Note that the high-mass end of the luminosity function in the near-IR is dominated by spheroids, or the progenitors of today's spheroids. If the use of the TP-AGB in the semi-analytic model shifts the luminosity function by $\\sim$1 mag at the high-mass end, it means that the mass-to-light ratio is lower by a factor of $\\sim$2.5 for a given luminosity. When galaxy masses are inferred from observations by the use of these models, they are lower by the same factor (as shown in M05). This may rise the question of whether there is enough mass in spheroids at high redshift to account for the $\\sim 50\\%$ of stellar mass in ellipticals measured in the local universe. However, the model correctly predicts the stellar mass density at all redshifts, meaning that only the \\textit{distribution of galaxy masses} is at tension with observations, \\textit{if the TP-AGB is not taken into account}. In fact, hierarchical models in general predict a faster evolution of the high-mass end of the stellar mass function than currently inferred from observations (see for instance Conselice et al. 2007). A more accurare derivation of galaxy masses through complete stellar population models with the TP-AGB, coupled with more accurate predictions from hierarchical models with the right input SSP, surely contribute to alleviate the discrepancy.  The inclusion of the TP-AGB allows the semi-analytic model to reproduce the very red end of the galaxy population at $z \\sim 2$, both for nearly-passive and for star-forming objects. It allows the model to do so with a comfortable range of galaxy masses and dust reddening. Most importantly, it contributes to a realistic and comprehensive treatment of the galaxy light emission in galaxy formation models, making them a much more precise tool to test our understanding of galaxy assembly.  The implementation of the TP-AGB allows the model to produce, at a given stellar mass, redder and more luminous galaxies in the near-IR, expecially at high redshift where the ages of the stellar populations peak around the epoch of maximal emission from this stellar phase. In case of nearly-passively evolving galaxies, the model can reproduce the red colours and high K-band magnitudes without invoking too large stellar masses or too old ages, which would be problematic in the hierarchical context. In the case of star-forming galaxies, the TP-AGB still increases the near-IR luminosity and makes the galaxies redder, without offsetting the blue optical colours. Thus, observed red colours in the near-IR do not necessarily imply old ages and passive evolution, a fact that again would be problematic for the hierarchical picture at high redshift. In general, the introduction of the TP-AGB in the models is a step forward in reconciling the hierarchical assembly mechanism with the observations of the high-redshift universe."
    },
    "0910/0910.2821_arXiv.txt": {
        "abstract": "We present a sensitive 870\\,\\micron\\ survey of the Extended Chandra Deep Field South (ECDFS) combining 310 hours of observing time with the Large Apex BOlometer Camera (LABOCA) on the APEX telescope. The LABOCA ECDFS Submillimetre Survey (LESS) covers the full $30'\\times30'$ field size of the ECDFS and has a uniform noise level of $\\sigma_{870\\mu{\\rm m}}\\approx1.2$\\,mJy\\,beam$^{-1}\\,$. LESS is thus the largest contiguous deep submillimetre survey undertaken to date. The noise properties of our map show clear evidence that we are beginning to be affected by confusion noise. We present a catalog of 126 submillimetre galaxies (SMGs) detected with a significance level above 3.7\\,$\\sigma$, at which level we expect 5 false detections given our map area of 1260 arcmin$^2$.\\\\ The ECDFS exhibits a deficit of bright SMGs relative to previously studied blank fields but not of normal star-forming galaxies that dominate the extragalactic background light (EBL). This is in line with the underdensities observed for optically defined high redshift source populations in the ECDFS (BzKs, DRGs, optically bright AGN and massive K-band selected galaxies). The differential source counts in the full field are well described by a power law with a slope of $\\alpha=-3.2$, comparable to the results from other fields.  We show that the shape of the source counts is not uniform across the field. Instead, it steepens in regions with low SMG density. Towards the highest overdensities we measure a source-count shape consistent with previous surveys. The integrated 870\\micron\\ flux densities of our source-count models down to $S_{870\\mu{\\rm m}}=0.5$\\,mJy account for $>65\\%$ of the estimated EBL from {\\it COBE} measurements.  We have investigated the clustering of SMGs in the ECDFS by means of a two-point correlation function and find evidence for strong clustering on angular scales $<1'$ with a significance of $3.4\\,\\sigma$.  Assuming a power law dependence for the correlation function and a typical redshift distribution for the SMGs we derive a characteristic angular clustering scale of $\\theta_0=14''\\pm7''$ and a spatial correlation length of $r_0=13\\pm6\\,h^{-1}$\\,Mpc. ",
        "introduction": "One of the most significant findings of the IRAS survey was the identification of a population of ultraluminous infrared galaxies (ULIRGs) that emit the bulk of their bolometric luminosity at far-IR wavelengths \\citep{sanders96}.  Surveys in the submillimetre and millimetre wavebands over the past decade have shown that ULIRGs are much more common at high redshift compared to the local universe \\citep[e.g.][]{barger99,cowie02,borys03,webb03,greve04,laurent05, coppin06,pope06,bertoldi07,beelen08,knudsen08,scott08,austermann09}. These surveys show that the comoving volume density of luminous submillimetre galaxies (SMGs) increases by a factor of $1000$ out to $z\\sim 2$ \\citep{chapman05}. Therefore luminous obscured galaxies at high redshift could dominate the total bolometric emission of galaxies at those epochs \\citep{blain99,lefloch05}   The identification and study of submillimetre galaxies has proved challenging since their first detection.  The limited mapping speed of typical (sub)millimetre bolometer cameras meant that only few very bright examples have been found, although gravitational lensing initially aided somewhat \\citep[e.g.][]{smail97,ivison98}.  Attempts to map large fields at submillimetre wavelengths have involved the use of patchworks of small ''jiggle'' maps \\citep[e.g.][]{coppin06} or mixtures of single-bolometer photometry, small jiggle maps and shallow scan maps used to construct a ''Super-map'' of GOODS-N \\citep{borys03,pope06}. Both of these approaches raise concerns about the homogeneity of the resulting maps and hence the reliability of the resulting source catalogues.  Scan maps, where the array is continuously moved on the sky to trace out a closed pattern, should result in much more homogeneous coverage and mapping, while at the same time allowing for a reliable removal of the bright emission from the atmosphere.  This technique has been used at submillimetre and millimetre wavelengths \\citep[e.g. at 350$\\mu$m,][at 1100$\\mu$m]{kovacs06,austermann09}, however, no deep survey have employed such a technique in the 870-$\\mu$m window where most of the published work on SMGs has been undertaken. Drawing this distinction between 870-$\\mu$m and 1100-$\\mu$m surveys may appear surprising given the modest difference between the two wavelengths and the assumed unstructured nature of the dust spectrum at these wavelengths. Despite only a 25\\% difference in the two wavelengths, there are hints of significant differences in the populations identified at 870-$\\mu$m and 1100-$\\mu$m \\citep[e.g.][]{greve04,younger08}, although these may in turn reflect the different mapping techniques used in individual studies.   The advent of the new Large APEX Bolometer Camera \\citep[LABOCA,][]{siringo09}, with an instantaneous 11.4$'$ field of view, on the 12-m APEX telescope \\citep{guesten06} provided the opportunity to undertake the first sensitive and uniform panoramic survey of the extragalactic sky at 870$\\mu$m.  To exploit this opportunity a number of groups within the Max Planck Gesellschaft (MPG) and the European Southern Observatory (ESO) communities proposed a joint public legacy survey of the Extended {\\it Chandra} Deep Field South (hereafter ECDFS) to the MPG and ESO time allocation committees: the LABOCA ECDFS submillimetre survey (LESS). The ECDFS covers a 0.5$^\\circ\\times$\\,0.5$^\\circ$ region centered on the {\\it Chandra} Deep Field South (CDFS) at RA $03^h32^m28^s.0$ Dec $-27^{\\circ}48^{'}30^{''}.0$. This field has very low far-infrared backgrounds and good ALMA visibility and hence has become one of the pre-eminent fields for cosmological survey science. As a result, the ECDFS is unique in the Southern Hemisphere in the combination of area, depth and spatial resolution of its multiwavelength coverage from X-rays through optical, near- and mid-infrared to the far-infrared and radio regimes. The central part of this field is coincident with the CDFS \\citep{giacconi02} which has now reached a depth of 2\\,Ms \\citep{luo08} and the deep {\\it Hubble Space Telescope} ({\\it HST}) imaging of the GOODS-S field \\citep{giovalisco04} and the {\\it Hubble} Ultra Deep Field \\citep[UDF,][]{beckwith06}.  In addition to the extremely deep observations of the central regions of this field as part of the CDFS, GOODS and {\\it Hubble} Ultra Deep Field surveys, the full 0.5$^\\circ$ field has extensive multiwavelength imaging available including: 250-ks {\\it Chandra} integrations over the whole field \\citep{lehmer05}; deep and multi-band optical imaging by COMBO-17 \\citep[][2008]{wolf04} and MUSYC \\citep{gawiser06} including {\\it HST} imaging for the GEMS project \\citep{caldwell08}; near-infrared imaging by MUSYC \\citep{taylor09}; deep mid-infrared imaging with IRAC as part of SIMPLE \\citep{damen09} and using the MIPS instrument at 24, 70 and $160\\mu$m by FIDEL (Dickinson \\etal\\ in prep.).  Longer wavelength coverage comes from BLAST \\citep{devlin09} at 250, 350 and $500\\mu$m (and in the near future from {\\it Herschel}), while radio coverage of this field is reported by \\citet[][]{miller08} and \\citet[][]{ivison09}.   The LABOCA survey of the ECDFS adds a waveband that pin-points the thermal emission from luminous dusty galaxies at $z\\sim1$--8: a powerful addition to this singularly well-studied region -- ideally placed for VLT observations and early science follow-up with ALMA. The completed LESS project provides a representative, homogeneous and statistically-reliable sample of the SMGs with the high-quality, multiwavelength data required to yield identifications, constrain their redshifts, bolometric luminosities and power sources and hence determine their contribution to the total star formation density at high redshift. These sources can be related in unprecedented precision to other populations of AGN and galaxies within the same volume to understand the place of SMGs in the formation and evolution of massive galaxies at high redshift. The survey is also sufficiently large that it should also yield examples of rare classes of SMGs, such as very high redshift sources, $z>4$ \\citep{coppin09}. These same data also provide submillimetre coverage of large numbers of high-redshift galaxies and AGN to determine their bulk submillimetre properties from the stacking analysis of sub-samples as a function of population, redshift, environment, etc. \\citep{greve09,lutz09}.  Together, these two techniques allow us to sample two orders of magnitude in bolometric luminosity -- from hyperluminous infrared galaxies with $10^{13}$\\,L$_\\odot$ which are directly detected in the LABOCA maps, down to luminous infrared galaxies at $10^{11}$\\,L$_\\odot$ which are detected statistically through stacking.  This range in luminosity encompasses the variety of populations expected to dominate the bolometric emission at $z \\sim 1$--3 and the cosmic submillimetre background.  In this paper we present a detailed description of the observations, reduction and analysis of the LABOCA observations of the ECDFS and the resulting catalogue of submillimetre galaxies.  The observations are described in \\S2, \\S3 presents our results and we discuss these in \\S4. Finally, in \\S5, we give our summary and the main conclusions of this work. We assume a cosmology with $H_0=70$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_\\Lambda=0.7$ and $\\Omega_M=0.3$. ",
        "conclusions": "We have presented a deep 870\\,\\micron\\ survey of the ECDFS using LABOCA on the APEX telescope at Llano de Chajnantor in Chile. This is the largest contiguous deep submm survey to date. Our map has a highly uniform noise level across the full $30'\\times30'$ field of 1.2\\,mJy$\\,{\\rm beam}^{-1}$ and our survey is $>95\\%$ complete for sources down to a flux limit of 6.5\\,mJy. Our main findings are summarized as follows:  \\begin{itemize}  \\item At the (beam smoothed) spatial resolution of $27''$ of our survey we find that the map's noise level is affected by confusion noise arising from faint, individually undetected SMGs. From the rms noise as a function of integration time we derive a confusion noise of $\\sigma_{\\rm c}\\approx\\, 0.9\\,{\\rm mJy}\\,{\\rm beam}^{-1}$.  \\item We identify 126 submm sources in a search area of 1260\\,arcmin$^2$ above a signal to noise threshold of $3.7\\,\\sigma$, which corresponds to an expected false detection rate of 5 sources.  \\item We have determined the differential and integrated source counts using a $P(D)$ analysis and an estimate based on our source catalog. Both results are in reasonable agreement and show that SMGs in the ECDFS are underabundant by a factor of $\\sim2$ for sources brighter than 3\\,mJy compared to the average of previous surveys. Under the assumption that the bulk of the sources are at $z>0.5$, this implies an underdensity of ULIRGs with $\\lfir>2\\times10^{12}\\lsol$ compared to other blank fields that have been observed in the submm. The source counts are well described by a single power law with a slope of $\\alpha=3.2\\pm0.2$.  \\item We derive the angular two-point correlation function for the SMGs and find clustering on angular scales $<1'$ with a significance up to $3.4\\,\\sigma$. Assuming a power law dependence for the correlation function we derive a clustering amplitude of $A_w=0.011\\pm0.0046$ or a characteristic angular scale of $\\theta_0=14''\\pm7''$ for $\\gamma =1.8$. Assuming a redshift distribution similar to that observed for spectroscopically confirmed SMGs, we derive a correlation length of $r_0=13\\pm6\\,h^{-1}$\\,Mpc, somewhat larger than previous estimates of the 3-D clustering of SMGs but in agreement with the clustering derived for 24\\,\\micron\\ selected ULIRGs.  \\item We have investigated for the first time the spatial variations of the SMG source counts. We find that the differential source counts in regions with an overdensity of SMGs have a different shape compared to those with underdensities.  While the counts in underdense regions are well fitted by a single power law with a slope of $\\alpha=3.6\\pm0.3$, the counts in the overdensities are significantly shallower with $\\alpha=2.9\\pm0.2$. The counts in the overdensities are slightly better described by a broken power law or a Schechter function. For flux densities below 8\\,mJy we find a slope of $\\alpha=2.4\\pm0.15$, for sources above this limit the counts are much steeper with $\\beta=4.7\\pm0.6$. This may indicate an intrinsic turn-over in the underlying luminosity function placing an upper limit on the FIR luminosity.  \\item The integrated 870\\,\\micron\\ flux density derived from our survey is $>29-32$\\,Jy\\,deg$^{-2}$ for sources brighter than $\\sim0.5$\\,mJy which corresponds to $>65-70$\\% of the extragalactic background light estimated from {\\it COBE} measurements. We do not find a significant difference of the quantity between SMG over- and underdensities. We conclude that ECDFS is underabundant of ULIRGs but not of more typical star forming systems with lower FIR luminosities, which dominate the extragalactic background light.  \\end{itemize}"
    },
    "0910/0910.1539_arXiv.txt": {
        "abstract": "We show that simple models of scalar-field dark energy leave a generic enhancement in the weak-lensing power spectrum when compared to the $\\Lambda$CDM prediction. In particular, we calculate the linear-scale enhancement in the convergence (or cosmic-shear) power spectrum for two best-fit models of scalar-field dark energy, namely, the Ratra-Peebles and SUGRA-type quintessence. Our calculations are based on linear perturbation theory, using gauge-invariant variables with carefully defined adiabatic initial conditions. We find that geometric effects enhance the lensing power spectrum on a broad range of scales, whilst the clustering of dark energy gives rise to additional power on large scales. The dark-energy power spectrum for these models are also explicitly obtained. On degree scales, the total enhancement may be as large as $30$-$40\\%$ for sources at redshift $\\sim1$. We argue that there are realistic prospects for detecting such an enhancement using the next generation of large telescopes. ",
        "introduction": "Over the past years, a concordance picture of our universe has emerged from a number of cosmological probes. In this paradigm, the dominant form of matter is cold and dark (CDM), but most of the mass-energy budget is in the form of dark energy, manifest in the accelerated expansion of the universe [for some of the earliest evidence from  Type Ia supernovae, see for example  \\cite{GoldPerl1998} and \\cite{Riess1998Aj}]. Some of the key goals in cosmology from the analysis of upcoming high-precision experiments within the next couple of decades are to determine the nature of dark energy, and whether there is a tension with its interpretation as a cosmological constant.  A decisive way to distinguish whether dark energy is the cosmological constant or some other dynamical entity is to establish whether the dark-energy equation-of-state parameter $w\\sub{DE}$ is constant or varies with redshift. The variation of $w\\sub{DE}$ can be probed via a range of astronomical techniques, including measurements of light-curves of type Ia supernova (SNIa), cosmic microwave background (CMB) anisotropies, galaxy redshift surveys, galaxy cluster abundance and cosmological weak lensing  (see \\cite{copeland,frieman+} for a summary of the various techniques).  Out of these techniques, weak lensing holds great promise in the understanding of dark energy, and has been identified as the most powerful individual technique, or most important component in a multi-technique study, provided systematic errors are well understood \\citep{detf, peacock}. Its utility as a dark-energy probe comes from sensitivity to both the expansion history of the universe and to the rate of growth of structures.  Observation of weak lensing is still in its infancy, with the first detection of cosmic shear made only nine years ago \\citep{wittman}. At present, there are a number of upcoming ambitious experiments - both terrestrial [\\eg LSST, ELT \\citep{lsst, elt}] and space-based [\\eg JDEM\\footnote{http://jdem.gsfc.nasa.gov}, Euclid\\footnote{http://sci.esa.int/euclid}], aimed at measuring cosmic shear to high accuracy. Given this rapid development, it is worth investigating if the upcoming weak-lensing experiments could discriminate simple models of dynamical dark energy from the cosmological constant.  Many previous works on this problem concentrate solely on the geometrical effects of dynamical dark energy. Furthermore, many authors assume dynamical dark energy with $w\\sub{DE}=$ constant, or some algebraic expression [numerous examples may be found the review by \\cite{copeland}]. Whilst this is a useful approximation for studying cosmological dynamics at late-time, it is impossible to extrapolate these simple forms of $w\\sub{DE}$ to early times, unless some \\ii{ad hoc} mechanisms are invoked. As a result, one cannot consistently analyse the perturbations in dark energy, and many authors then neglect these perturbations altogether. Even those who do include dark-energy perturbations often gloss over the issue of initial conditions, and simply assume that the perturbations are insensitive to initial conditions without any justification.  In this work, we shall calculate the effects of dynamical dark energy on weak-lensing measurements without making any of the assumptions in the previous paragraph. We shall analyse simple models of dynamical dark energy using gauge-invariant perturbation theory, with well-defined initial conditions. As we shall see, these models leave some interesting imprints on the weak-lensing observables on all scales. Ultimately, we shall give a basic assessment of the prospects for distinguishing dynamical dark energy from the cosmological constant using weak-lensing measurements.  One particularly important aspect of this work is the inclusion of dark-energy perturbations throughout our analysis. There have been a number of recent works examining the effects of dark-energy perturbations on cosmological observables \\citep{dent, sapone,hwang2}, including some interesting numerical simulations \\citep{alimi, jennings}. Our approach is complementary to these works, and goes further in quantitatively establishing the effects of dynamical dark energy on weak-lensing power spectra. As a by-product, we obtain explicitly, for the first time, the dark-energy power spectrum for these models. Such a spectrum is useful as it quantifies the clustering of dark energy as a function of physical length scale.  Throughout this paper, we work in Planck units in which $c=\\hbar=1$. ",
        "conclusions": "In this work, we have investigated whether there are observable signatures of dynamical dark energy in weak-lensing measurements. Our basic assumptions are that the background cosmology is a flat FRW universe, containing radiation, dark matter and the simplest type of scalar-field dark energy (with no hot dark matter, baryons or tensor modes). Our technique involves calculating the linear density power spectra by integrating six coupled differential equations, expressed in terms of gauge-invariant density and velocity variables for dark matter, radiation and dark energy (equations \\ref{d1}-\\ref{v3}). We assume adiabatic initial conditions for all three types of perturbations, and normalise the matter power spectrum (with primordial form $\\sim k^{n_s}$) on CMB scales. We base our analysis specifically on two simple models of quintessence, namely the Ratra-Peebles and SUGRA-type potentials, with model parameters that best fit the recent CMB and combined supernova data.  We find that the best-fit quintessence models yield matter power spectra that are hardly distinguishable from that of the $\\Lambda$CDM model (figure \\ref{figpm}) except on large scales where dark-energy clustering is strongest, as shown in figure \\ref{figde}. We believe this is the first time that the dark-energy power spectra for these oft-cited models are explicitly calculated.  When the weak-lensing signals are extracted from the 3D spectra, we found an imprint in the form of an all-scale enhancement in the convergence (or the cosmic shear) power spectrum compared with $\\Lambda$CDM cosmology. For sources at redshift 0.5, for instance, the enhancement may be as large as $40\\%$ at about $2^\\circ$ scale (or 20\\% at $z=1$). These enhancements can be attributed to two main reasons, namely i) the change in distance-redshift relation (figure \\ref{figdistance}), and ii) the lensing by large-scale dark-energy clusters (figure \\ref{figignore}).  Assuming that the primary sources of errors in the weak-lensing measurements will be from cosmic variance and intrinsic ellipticity dispersion in the shapes of distant galaxies, there are good  prospects for detecting this enhancement for sources at redshift $z\\sim0.5$ using future facilities with surveys covering most of the sky, and capable of measuring ellipticities for hundreds of galaxies per arcmin$^{2}$ (figure \\ref{figerrors}). On the timescale of a decade, one class of telescope capable of such a survey is the E-ELT, 42m in diameter and including adaptive optics, for which operations are planned to start in 2018. To put this requirement in context, current ground-based surveys typically measure ellipticities for a few tens of galaxies per arcmin$^{2}$ over order 100 square degrees (e.g. CFHTLS-Wide\\footnote{http://terapix.iap.fr}), and the highest specification near-future survey DES\\footnote{http://www.darkenergysurvey.org} will measure a comparable number density of galaxies over 5000 square degrees. Here we restricted our consideration to distant galaxies at the same redshift, but for a real survey with sources distributed in redshift, photometric redshift estimates can be obtained using observations in multiple filters. Tomographic information using the auto- and cross-correlations between the cosmic shear in various redshift bins, will give additional leverage in distinguishing between models of dark energy.  \\bbb  \\no{\\bf"
    },
    "0910/0910.3827_arXiv.txt": {
        "abstract": "{Counting clusters is one of the methods to constrain cosmological parameters, but has been up to now limited both by the redshift range and by the relatively small sizes of the homogeneously surveyed areas.} {In order to enlarge publicly available optical cluster catalogs, in particular at high redshift, we have performed a systematic search for clusters of galaxies in the Canada France Hawaii Telescope Legacy Survey (CFHTLS).} {We considered the deep 2, 3 and 4 CFHTLS Deep fields (each 1$\\times$1~deg$^2$), as well as the wide 1, 3 and 4 CFHTLS Wide fields. We used the Le Phare photometric redshifts for the galaxies detected in these fields with magnitude limits of i'=25 and 23 for the Deep and Wide fields respectively. We then constructed galaxy density maps in photometric redshift bins of 0.1 based on an adaptive kernel technique and detected structures with SExtractor at various detection levels. In order to assess the validity of our cluster detection rates, we applied a similar procedure to galaxies in Millennium simulations. We measured the correlation function of our cluster candidates. We analyzed large scale properties and substructures, including filaments, by applying a minimal spanning tree algorithm both to our data and to the Millennium simulations. } {We have detected 1200 candidate clusters with various masses (minimal masses between 1.0 10$^{13}$ and 5.5 10$^{13}$ and mean masses between 1.3 10$^{14}$ and 12.6 10$^{14}$ M$_\\odot$) in the CFHTLS Deep and Wide fields, thus notably increasing the number of known high redshift cluster candidates. We found a correlation function for these objects comparable to that obtained for high redshift cluster surveys.  We also show that the CFHTLS deep survey is able to trace the large scale structure of the universe up to z$\\geq$1.  Our detections are fully consistent with those made in various CFHTLS analyses with other methods. We now need accurate mass determinations of these structures to constrain cosmological parameters.} {We have shown that a search for galaxy clusters based on density maps built from galaxy catalogs in photometric redshift bins is successful and gives results comparable to or better than those obtained with other methods. By applying this technique to the CFHTLS survey we have increased the number of known optical high redshift cluster candidates by a large factor, an important step towards using cluster counts to measure cosmological parameters.} ",
        "introduction": "\\label{sec:intro}  The beginning of the 21st century is an exciting period for cosmological studies. Several methods now allow to put strong constraints on cosmological parameters. We can for example reconstruct Hubble diagrams (supernovae or tomography) or use directly the primordial fluctuation spectrum.  In addition, the cluster count technique is probably the oldest one (see e.g. Gioia et al., 1990). Up to now this technique was penalized by the redshift range of detected clusters, which was too low to make the difference between flat and open universes. Distant cluster surveys have also been mainly conducted in areas too small or with inhomogeneous selection functions.  Besides cluster mass knowledge, this technique requires indeed large fields of view of several dozen square degrees to provide large numbers of cluster detections at z$\\geq$1 (e.g. Romer et al. 2001).  Recent X-ray cluster surveys are beginning to produce cluster catalogs at high z (e.g. the XMM-LSS survey, Pierre et al., 2007) and it is the goal of the present paper to contribute to the production of similar large cluster catalogs based on optical Canada France Hawaii Telescope Legacy Survey data.  The Canada-France-Hawaii Telescope Legacy Deep and Wide Surveys (CFHTLS-D and CFHTLS-W) respectively explore solid angles of 4 deg$^2$ and 171 deg$^2$ of the deep Universe, each in 4 independent patches (http://www.cfht.hawaii.edu/Science/CFHLS/). For both surveys, observations are carried out in five filters ($u*,g',r',i'$ and $z'$) providing catalogs of sources that are 80\\% complete up to $i_{AB}$=26.0 (CFHTLS-D) and $i_{AB}$=24.0 (CFHTLS-W) (Mellier et al 2008, http://terapix.iap.fr/cplt/oldSite/Descart/CFHTLS-T0005-Release.pdf). The CFHTLS-W, in particular, encloses a sample of about 20 10$^6$ galaxies inside a volume size of $\\sim 1$ Gpc$^3$, with a median redshift of z$ \\sim 0.92$ (Coupon et al 2009). According to the standard cosmological model, the CFHTLS-W (W1, W2, W3, and W4 herafter) is then expected to contain 1000 to 5000 clusters of galaxies with accurate photometric redshifts, most of them in the $0.6<z<1.5$ range. Likewise, the CFHTLS-D (D1, D2, D3, and D4 hereafter) should contain 50 to 200 clusters, with a significant fraction at higher redshift than the CFHTLS-W.  The two surveys are therefore complementary data sets. They can produce together two homogeneous optically selected samples of clusters of galaxies that can be reliably compared and used jointly to explore the mass function, the abundance of clusters of galaxies and the evolution of cluster galaxy populations as a function of lookback time.  The construction of homogeneous catalogs of optically selected clusters of galaxies is not, a simple task. Early searches for clusters of galaxies in the CFHTLS were performed by Olsen et al. (2007) based on a matched filter detection algorithm applied to the Deep fields (see also the recent paper by Grove et al. 2009). Galaxy density maps combined with photometric redshift catalogs were considered by Mazure et al. (2007) in the D1 field.  Lensing techniques were also employed to detect massive structures in the CFHTLS (e.g. Cabanac et al. 2007, Gavazzi $\\&$ Soucail 2007, Berg\\'e et al. 2008). Other cluster studies based on the CFHTLS data (e.g. the CFHTLS-CARS survey: Erben et al. 2009, Hildebrandt et al. 2009) and based for part of them on the red sequence in the color magnitude diagram are also in progress.  We present here a systematic search for galaxy clusters in the D2, D3 and D4 Deep fields (the D1 field was already analyzed by Mazure et al. 2007), as well as in the regions of the W1, W3 and W4 Wide fields available in the T0004 release. Our approach is based on photometric redshifts computed for all the galaxies extracted in each field (Coupon et al. 2009). In this way, we take into account the full color information and not only two bands as for example in red sequence searches.  We divided the galaxy catalogs in slices of 0.1 in redshift, each slice overlapping the previous one by 0.05, and built density maps for each redshift slice. Structures in these density maps were then detected with the SExtractor software in the different redshift bins. We applied the same method to similar size mock samples built from the Millennium simulation, in order to estimate the reliability of our detections. We measured the clustering properties of our catalog. We then analyzed substructuring and filamentary large scale properties by applying a minimum spanning tree algorithm both to our data and to the Millennium simulation.  In this paper we assume H$_0$ = 70 km s$^{-1}$ Mpc$^{-1}$, $\\Omega _m$=0.3, $\\Omega _{\\Lambda}$=0.7. All coordinates are given at the J2000 equinox and magnitudes are in the AB system. ",
        "conclusions": "We have detected 1200 candidate clusters in the CFHTLS Deep and Wide fields. Statistically, more than 80$\\%$ are real structures at z$\\leq$1. This is confirmed by internal and external comparisons with literature catalogs.  \\begin{table} \\caption{Number and density of $detections$ as a function of S/N.} \\begin{center} \\begin{tabular}{ccc} \\hline SExtractor threshold & number of $detections$ & density of $detections$ \\\\ &  &  deg$^{-2}$ \\\\ \\hline 2 & 535 & 19.2 \\\\ 3 & 262 & 9.4 \\\\ 4 & 170 & 6.1 \\\\ 5 & 98 & 3.4 \\\\ 6 & 135 & 4.8 \\\\ \\hline \\end{tabular} \\end{center} \\label{tab:summary} \\end{table}  Table~\\ref{tab:summary} gives the number and density of $detections$ as a function of S/N. Over an effective area of $\\sim$28 deg$^2$ (see Coupon et al. 2009) this means that we detect 19.2 candidate clusters per deg$^2$  for mass $\\geq$ 1.0 10$^{13}$M$_\\odot$, 9.4 candidate clusters per deg$^2$  for mass $\\geq$ 1.3 10$^{13}$M$_\\odot$, 6.1 candidate clusters per deg$^2$  for mass $\\geq$ 3.3 10$^{13}$M$_\\odot$, 3.4 candidate clusters per deg$^2$  for mass $\\geq$ 3.5 10$^{13}$M$_\\odot$, and 4.8 candidate clusters per deg$^2$  for mass $\\geq$ 5.5 10$^{13}$M$_\\odot$.  Given the typical uncertainty on detection rates, the two last candidate cluster densities are compatible. We note that these numbers are not corrected for detection efficiency. These numbers are also fully compatible with recent X-ray estimates (e.g. Pacaud et al. 2007) showing detections of 5.8 clusters per deg$^2$ for masses greater than 4.1 10$^{13}$M$_\\odot$ (number scaled to our cosmology).  If we compare our results to published optically based cluster catalogs, our survey represents a major step forward (see Table~\\ref{tab:summary2}). We have one of the deepest and largest cluster catalogs in the CFHTLS area by a factor of $\\sim$10 most of the time. We also basically provide the only cluster detections at z$\\geq$1 using CFHTLS data. The comparison of our resuls to the well known MaxBCG SDSS catalog (Koester et al. 2007) provides similar numbers in terms of cluster spatial density. The present deep and wide surveys provides more than 13000 and 10500 $detections$ per Gpc$^3$. This is comparable to the 16735 clusters per Gpc$^3$ of Koester et al. (2007), assuming the values of Table~\\ref{tab:summary2}.  \\begin{table*} \\caption{Comparison of our cluster detections with public CFHTLS optically based cluster catalogs and with the MaxBCG SDSS catalog.} \\begin{center} \\begin{tabular}{rrrrr} \\hline Authors & Detection method & Covered area & Number of detections & Maximal redshift \\\\ &  & deg$^2$ &  &  \\\\ \\hline Present paper Deep & Photometric redshifts & 2.5 & 171 & 1.5 \\\\ Present paper Wide & Photometric redshifts & 28 & 1029 & 1.2 \\\\ Olsen et al. (2007) & Matched Filter & 4 & 162 & 1.15 \\\\ Mazure et al. (2007) & Photometric redshifts & 1 & 44 & 1.5 \\\\ Cabanac et al. (2007) & Strong lensing & 28 & 40 & 1. \\\\ Limousin et al. (2009) & Strong lensing & 102 & 13 & 0.85 \\\\ Gavazzi $\\&$ Soucail (2007) & Weak lensing & 4 & 14 & 0.55 \\\\ Berg\\'e et al. (2008) & Weak lensing & 4 & 7 & 0.5 \\\\ Koester et al. (2007) & Max BCG & 7500 & 13823 & 0.3 \\\\ \\hline \\end{tabular} \\end{center} \\label{tab:summary2} \\end{table*}    These results illustrate the power of optical deep and wide field surveys to provide large samples of galaxy clusters. These samples could be used for pure cosmological applications (e.g. based on cluster counts, Romer et al.  2001) or more generally for the study of structures within a broad mass range. In particular, it is remarkable that our deep catalogs are among the first ones to provide numerous group detections at redshifts greater than 1.  In this picture, the perspectives of our work are:  - To perform a more homogeneous comparison with matched filter detections across the Wide fields, and this will be the subject of a future paper. A by-product of this future work will be the center refinement of our candidate clusters.  - To define a more precise optically-based mass estimator, via e.g. the galaxy luminosity functions. This requires however to have performed the previous match.  - To assess more precisely the detection rate for very massive clusters. The Millennium simulation as it is now does not provide a large enough number of very massive structures and this produces too large an uncertainty. This is prohibitive for any serious cosmological application based on cluster counts, since the most massive clusters are the most constraining for cosmology (e.g. Romer et al. 2001) and detection rates for these massive clusters need to be precisely evaluated. This aspect is currently under work and will also be developed in future papers.  - Finally, the spectroscopic confirmation of the z$\\geq$1 candidate clusters we have detected would be crucial for cosmology.  The cluster list will be available via the Cencos database at http://cencosw.oamp.fr/ in a near future."
    },
    "0910/0910.4900.txt": {
        "abstract": "\\abstract % context heading (optional) % {} leave it empty if necessary {X ray clusters are conventionally divided into two classes: ``cool core'' (CC) and ``non cool core'' (NCC) objects, on the basis of the observational properties of their central regions. Recent results have shown that the cluster population is bimodal (Cavagnolo et al. 2009).} % aims heading (mandatory) {We want to understand whether the observed distribution of clusters is due to a primordial division into two distinct classes rather than to differences in how these systems evolve across cosmic time.} % methods heading (mandatory) {We systematically  search the ICM of NCC clusters in a subsample of the B55 flux limited sample of clusters for regions which have some characteristics typical of cool cores, namely low entropy gas and high metal abundance} % results heading (mandatory) {We find that most NCC clusters in our sample host regions reminiscent of CC, i.\\, e.\\, characterized by relative low entropy gas (albeit not as low as in CC systems) and a metal abundance excess. We have dubbed these structures ``cool core remnants'', since we interpret them as what remains of a cool core after a heating event (AGN giant outbursts in a few cases and more commonly mergers). We infer that most NCC clusters have undergone a cool core phase during their life. The fact that most cool core remnants are found in dynamically active objects provides strong support to  scenarios where cluster core properties are not fixed ``ab initio'' but evolve across cosmic time.} % conclusions heading (optional), leave it empty if necessary {} ",
        "introduction": "Galaxy clusters are often divided by X-ray astronomers into two classes: ``cool core''(CC) and ``non-cool core'' (NCC) clusters. The former are characterized by a set of properties: a prominent surface brightness peak, usually roughly coincident with the center of the large scale X-ray isophotes and with the position of the brightest central galaxy (BCG), associated to a decrease of the temperature profile in the inner regions and to a positive gradient of the metal abundance profile. In the central regions of these clusters, the cooling time is significantly smaller than the Hubble time but the radiative losses are balanced by some form of heating, that in the currently prevailing scenario is attributed to the central AGN, which prevents the gas from cooling indefinitely and  flowing inward. NCC clusters are characterized by the lack of these observational features. Several indicators based on these observational characteristics have been proposed to classify clusters into these categories: e.\\,g.\\, temperature drop (\\citealt{sanderson06, sanderson09}), cooling time \\citep{bauer05}, a composition of these two criteria (\\citealt{dunn05, dunn08}), the slope of the gas density profile at a given radius \\citep{vikh_garching} and core entropy \\citep{cava09}. In a recent paper (\\citealt{pap1}, hereafter Paper I), we suggested a robust indicator to classify clusters (see also Sect.\\, \\ref{sec_classif}), based on pseudo-entropy gradients. The classification based on this indicator improves the traditional classification scheme based on the temperature drop and is essentially equivalent to classifications based on the cooling time. Increasing attention has been devoted to the statistics of CC, both in the local Universe and at higher redshifts. While the precise results depend strongly on the indicator used to classify clusters, the fraction of CC is considered to be about $50\\%$ of the clusters population. There are some indication that  the cluster population is bimodal \\citep{cava09} but there are also some intermediate objects which are not easily classified (Paper I).\\\\ One of the open questions in the study of galaxy clusters, concerns the origin of this distribution. The original model which prevailed for a long time assumed that the CC state was a sort of ``natural state'' for the clusters and the observational features were explained within the context of the old ``cooling flow'' model: radiation losses cause the gas in the centers of these clusters to cool and to flow inward. Clusters were supposed to live in this state until disturbed by a ``merger''. Indeed, mergers are very energetic events that can shock heat \\citep{burns97} and mix the ICM \\citep{gomez02, ritchie02}: through these processes they were supposed to efficiently destroy cooling flows. After the mergers, clusters were supposed to relax and go back to the cooling flow state in a sort of cyclical evolution. With the fall of the ``cooling flow'' brought about by the \\xmmn and \\chandra observations (e.\\,g.\\, \\citealt{peterson01} and \\citealt{molepizzo01}), doubts were cast also on the interpretation of mergers as the dominant mechanism which could transform CC clusters into NCC. More generally speaking, the question arose whether the observed distribution of clusters was due to a primordial division into the two classes rather than to evolutionary differences during the history of the clusters.  \\\\ \\citet{mccarthy04, mccarthy08} noticed the absence of systems that resemble observed NCC clusters in cosmological simulations, suggesting that mergers could not be the origin of the cluster distribution. They proposed that early episodes of non gravitational pre-heating may  explain the dichotomy; they envisage a scenario in which NCC clusters have been pre-heated to levels greater than $\\sim 300\\, \\rm{keV}\\, \\rm{cm}^{2}$ and do not have enough time to develop a cool core while CC clusters have been pre-heated to lower levels and need an additional source of present-day heating to offset cooling. With this model, \\citet{mccarthy08} explained observed entropy and gas density \\chandra profiles for CC and \\emph{ROSAT} gas density profiles for NCC. \\citet{ohara06} supported the primordial model showing that the scatter in scaling relations is larger for CC clusters than for NCC, suggesting that CC are not more relaxed than NCC systems. Moreover the  two body idealized simulations by \\citet{poole08} showed that mergers cannot produce extended ``warm'' cores and cannot destroy metal abundance gradients. \\\\ However, the evolutionary ``merger'' scenario has been continuously supported by observations. For instance, \\citet{sanderson06} compared temperature and cooling time profiles of a sample of clusters observed by \\chandra with indicators of merger activity and suggested that merger may be the primary factor in preventing the formation of CC. More recently, \\citet{sanderson09b} have shown in a large sample of objects that the X-ray/BCG projected offset correlates with the gas density profile, which can be considered an indicator of the CC state. If the X-ray/BCG offset ``measures'' the dynamical state of the cluster, this result implies that the cool core strengths diminishes in more dynamically disturbed clusters. Another indicator of dynamical activity is the presence of extended radio halos on Mpc scales, whose origin is likely related to turbulent acceleration driven by mergers (\\citealt{ferrari08} for a recent review). None of the clusters which are found to host a radio halo (on Mpc scales) in a complete survey \\citep{venturi08, brunetti09} can be classified as a CC object. These relations are indeed expected if mergers can efficiently destroy cool cores.\\\\ An interesting model has been proposed by \\citet{motl04}, who suggested a double role for mergers in the cluster formation process: while they can shock heat the ICM they also efficiently mix the ICM and participate to the formation of CC by providing cool gas. By analyzing simulated temperature and surface brightness maps, they find that observational signatures of a cool core may disappear after mergers but cool gas remains in the systems at all time. Indeed mergers are complicated phenomena during which the ICM of the interacting objects is mixed. Also  \\citet{ascas06}, with different simulations with the purpose of studying the formation of cold fronts, showed that an interacting subcluster may donate its cool gas to the main cluster and that, during the merging processes, many hydrodynamic effects contribute to the mixing of the gas. \\\\ Recently, we found an unexpected and interesting result from the analysis of the metalicity profiles in a large sample of clusters (Paper I): some clusters that cannot be classified as CC show a metal abundance excess at their center, witout a significant temperature decline or an X-ray brightness excess. We suggested that at least some NCC clusters have spent part of their lives as CC objects and therefore there cannot be a complete primordial separation of the two classes of objects.\\\\ However, further analysis is needed on this result to use it to gain insight on the origin of the distribution of CC-NCC objects. We would like to better characterize these regions and to know if and to what extent they are common in the population of NCC objects. The results in Paper I, as well as those described above comparing observations and simulation, are based on  ``uni-dimensional'' properties (entropy and metal abundance profiles). Indeed, global properties and unidimensional profiles are easy to derive and it is possible to compare them quantitatively but they inevitably result in a loss of information. There is now a large number of thermodynamic maps in the literature, which show that spherical symmetry is not generally fulfilled in clusters and revealed the presence of regions with ``anomalous'' characteristics outside of the cores (e.g \\citealt{sun02,rossetti06} and many others ). \\\\ In this paper, we present a systematic two-dimensional analysis on a sample of 35 clusters observed with \\xmm,  to look for and characterize regions with a significant metal excess in NCC clusters, as those found in Paper I. The outline of the paper is as follows: in Sect.\\,2 we describe the sample and the data analysis, in Sect.\\,3 we discuss the classification schemes and we define ``CC remnants'' in Sect.\\,4.  We interpret our results in Sect.\\,5 and  we summarize our findings in Sect.\\,6. Quoted confidence intervals are 68\\% for one interesting parameter unless otherwise stated. All results are given assuming a $\\Lambda$CDM cosmology with $\\Omega_{\\rm{m}} = 0.3$, $\\Omega_\\Lambda = 0.7$, and $H_0 = 70$ km s$^{-1}$ Mpc$^{-1}$.   %__________________________________________________________________ ",
        "conclusions": "In this paper, we have presented a systematic two-dimensional analysis of the entropy and metal abundance of a sample of nearby clusters observed with \\xmmn. \\begin{itemize} \\item We have analyzed the observations of a sample of 35 clusters. Following the prescription in Paper I, we have classified 14 of them as low-entropy cores. In the remaining 21 non-LEC objects we have performed a systematical analysis of the entropy maps  and selected regions with pseudo-entropy ratio lower than $ 0.8$, for a proper spectral analysis.  All the clusters in our subsample host regions with these characteristics. \\item In most non-LEC objects (12/21) the low entropy regions are characterized by a significant metal abundance excess with respect to outer regions of galaxy clusters. For 6 of the remaining clusters, we can significantly exclude an abundance larger than $0.4Z_{\\sun}$, while for 3 objects the error bars are too large to discriminate. We have interpreted the low-entropy high metal abundance regions that we found in 12 clusters as what remains of a cool core after a heating event, and we have dubbed them ``CC remnants''. \\item For two clusters, A1650 and MKW3s, the most likely heating mechanism is AGN feedback. As shown by \\citet{voit05}, clusters with core entropy $\\lesssim 50\\, \\rm{keV}\\, \\rm{cm}^{2}$  are within the reach of powerful AGN outbursts. \\item In 8/12 clusters with CC remnant, the core entropy is too large to be produced by giant AGN outbursts. In all these clusters we have found indication of on-going interactions and 5 of them are probably undergoing a major merger. Therefore the most likely heating mechanism for this class of objects is merging, which can shock heat and mix the ICM. In most of these clusters, the CC remnant is associated to the BCG or to a giant elliptical galaxy which could be responsible for the production of the metal excess during the LEC phase. In A3667 and A2256 the merging may have completely decoupled the low-entropy metal rich gas from its associated galaxy concentration. In A754, A2256 and A3667 the position of the CC remnant does not coincide with the centroid of the X-ray emission, indicating that the low entropy metal rich ICM is not in equilibrium within the potential well of the host clusters. \\item In 6 objects out of 21, we have not found a significant metal excess in the low entropy regions. Since most of these clusters show indications of  major mergers, we have interpreted them as clusters where the metal abundance gradient has been erased by the mixing following mergers. \\end{itemize} The results presented in this paper strengthens those shown in Paper I. Indeed, in Paper I we found a metal abundance excess in a small fraction of non-LEC clusters. Thanks to the two-dimensional analysis presented in this paper, we have found a metal abundance excess in most objects. We conclude that most non-LEC objects have spent part of their life as low-entropy cores. After a heating event the LEC signature disappears but the ``CC-like'' ICM (i.\\,e.\\, characterized by large metal abundance and low entropy, albeit not as low as in LEC) remains in the systems and can be identified in two-dimensional entropy maps. In the framework of the alternative ``primordial'' scenario the low entropy metal rich regions could be interpreted as ``progenitors'' of cool cores, that are now evolving (slowly cooling and increasing their metalicity) to become LEC.  However this scenario does not explain why most of these regions lie in dynamically active objects and at least three of them are found in a configuration of non-equilibrium, offset from the center of the clusters. Conversely,  this issue is naturally addressed if they are the results of a merger event in the evolutionary scenario.\\\\ The results summarized above strongly support scenarios where cluster core properties are not fixed ``ab initio'' but evolve across cosmic time."
    },
    "0910/0910.1225_arXiv.txt": {
        "abstract": "{ The coupling between time-dependent, multidimensional MHD numerical codes and radiative line emission is of  utmost importance in the studies of the interplay between dynamical and radiative processes in many astrophysical environments, with particular interest for problems involving radiative shocks. There is a widespread consensus that line emitting knots observed in Herbig-Haro jets can be interpreted as radiative shocks. Velocity perturbations at the jet base steepen into shocks to emit the observed spectra. To derive the observable characteristics of the emitted spectra, such as line intensity ratios, one has to study physical processes that involve the solution of the MHD equations coupled with radiative cooling in non-equilibrium conditions. }{ In this paper we address two different aspects relevant to the time-dependent calculations of the line intensity ratios of forbidden transitions, resulting from the excitation by planar, time-dependent radiative shocks traveling in a stratified medium. The first one concerns the impact of the radiation and ionization processes included in the cooling model, and the second one the effects of the numerical grid resolution. }{ Dealing with both dynamical and radiative processes in the same numerical scheme means to treat phenomena characterized by different time and length scales. This may be especially arduous and computationally expensive when discontinuities are involved, such as in the case of shocks. Adaptive Mesh Refinement (AMR) methods have been introduced in order to alleviate these difficulties. In this paper we apply the AMR methodology to the treatment of radiating shocks and show how this method is able to vastly reduce the integration time. }{ The technique is applied to the knots of the HH 30 jet to obtain the observed line intensity ratios and derive the physical parameters, such as density, temperature and ionization fraction. We consider the impact of two different cooling functions and different grid resolutions on the results. }{ We conclude that the use of different cooling routines has effects on results whose weight depends upon the line ratio considered. Moreover, we find the minimum numerical resolution of the simulation grid behind the shock to achieve convergence in the results. This is crucial for the forthcoming 2D calculations of radiative shocks. } ",
        "introduction": "Supersonic flows are ubiquitous in the Universe: expanding supernova remnants, stellar winds, AGN and Herbig-Haro jets are some examples, and shocks are abundant and prominent in these flows. Shocks located in extragalactic environments, such as AGN jets, can be considered adiabatic, since the cooling time for thermal emission typically exceeds  by far the source lifetime. On the other hand, shocks that form in galactic objects such as supernova remnants and Herbig-Haro jets must be treated including radiative effects.  Radiative shocks have been studied in steady-state conditions by several authors (e.g. Cox \\& Raymond \\cite{CR85}, Hartigan et al. \\cite{HA94}), who derived the one-dimensional post-shock behavior of various physical parameters (temperature, ionization fraction, electron density, etc.), as functions of the distance from the shock front. From these studies it turns out that the physical parameters vary behind the shock front on scale lengths that differ by orders of magnitude, and this becomes a major problem when tackling the time-dependent evolution of such radiative shocks. For example, in Herbig-Haro jets, one needs to treat in the same scheme spatial scales well below $\\sim 10^{13}$ cm, to reproduce the time-dependent post-shock temperature variations correctly, and scales $\\lapp \\ 10^{15}$ cm to study the behavior of the electron density (Massaglia et al. \\cite{MA05a}, Paper I). Under these conditions and with these requirements, employing an Adaptive Mesh Refinement (AMR) technique becomes almost mandatory. This technique can provide adequate spatial and temporal resolution by dynamically adapting the grid to the moving regions of interest, giving a tremendous saving in computational time and memory overhead with respect to the more traditional uniform grid approach.  The important issue of numerical spatial resolution in time-dependent simulations of radiative shocks in 2D has been considered by \\cite{Raga07} employing an AMR technique as well. They discussed the dependence of the morphological structure of the perturbations depending on the grid resolution. Their maximum grid resolution was with a cell size corresponding to $1.5 \\times 10^{12}$ cm. They found that while the detailed structure of the shocks depends on the resolution, the emission line luminosities, integrated over the volume, are less dependent on cell size.  A second crucial question is: does the dependence on the cooling function details have qualitative or quantitative effects on the calculated distributions of the physical parameters? Te\\c{s}ileanu et al. (\\cite{TA08}) have developed a detailed treatment including an ionization network of 29 species and compared, at constant numerical resolution, the resulting distribution of physical parameters (such as temperature, density, etc.) with the one obtained employing the simplified cooling of Paper I, for a 2D cylindrical jet affected by perturbations that evolved in internal working surfaces. They found quantitative differences but qualitatively the outcomes were very similar.  A more challenging problem than the determination of the integrated line emissions and distributions of physical parameters is the calculation of line intensity ratios between different species. In this case, the fractional abundances of the various species vary in very different ways with space behind the shocks (see Paper I). This variation is mainly governed by the temperature profile, which is typically a very steep function of space. Therefore we will examine the effects of both the details of the cooling function and of the grid resolution adopted.  As discussed in Paper I and Massaglia et al. (\\cite{MA05b}), observations of some HH jets at distances of a few arcseconds from the source typically show that the behavior of temperature, ionization and density along the jet is incompatible with a freely cooling jet. It is now generally accepted that the line emission in HH jets is a result of the observation of several post-shock regions of high excitation with a filling factor of about 1\\%. In our calculations, we will refer, in particular, to the observational results described in Bacciotti et al. (\\cite{BE99}), where a specialized diagnostic technique has been applied to HST data of the HH 30 jet. Hartigan \\& Morse (\\cite{HM07}) also re-examined the HH 30 case using slitless spectroscopy performed with the Hubble Space Telescope Imaging Spectrograph. Their results are fully consistent with the findings in Bacciotti et al. (\\cite{BE99}).  In this paper, we will follow the dynamical evolution of an initial perturbation as it steepens into a (radiative) shock traveling along the jet, and derive the post-shock physical parameters consistently. From these parameters, we construct the synthetic theoretical emission line ratios to be compared with observations. Using this setup, we implicitly consider that the medium in the real jet returns close to the steady jet conditions between the propagating shocks. This was verified in parallel 2D simulations (Te\\c{s}ileanu et al. \\cite{TA09}), as the 1D simulation, being done in the reference frame of the mean flow, lacks the steady flow of the jet. After the shock and the high-density post-shock zone pass, there follows an underdense zone, and then the medium returns to the stationary conditions of the constant flow. The distance after which this happens increases during the evolution of the shock from $0.5\\times 10^{14}$cm close to the jet base to about $10^{15}$cm at 500AU from the base. This increasing distance becomes resolved observationally at about 200AU from the base of the jet, and also at such distances interactions between different shockwaves become likely, these being possible explanations for the oscillations visible in the observational data after 200AU. The emission at each point of the jet is computed by simulating the propagation of a shock produced at the base of the jet up to that point. As discussed before, we will adopt two different treatments for the radiative losses, namely: i) the simplified cooling employed in Paper I and ii) the new cooling treatment described in Te\\c{s}ileanu et al. (\\cite{TA08}).  The plan of the paper is the following: In Section 2 we present the computation setup and the adopted techniques to model the problem; in Section 3 we apply the model to the case of HH 30 and discuss the results obtained with the two cooling functions and the effects of the grid resolution; the conclusions are drawn in Section 4. ",
        "conclusions": "We have demonstrated in previous works that numerical integration of time-dependent radiating (shocked) flows can substantially benefit from the employment of adaptive mesh refinement methods, and the present work supports this conclusion. By computing the cooling time scale according to the algorithm described in Mignone et al. (\\cite{MA09}), one could take full advantage of the time adaptation process. The use of this method allows us to tackle simulations of radiating shocks in 2D as they propagate along a cylindrical supersonic jet.  Our results applied to the DG Tau jet (Paper I) as well as the HH 30 jet (present paper) revealed a good agreement between the observed and calculated line intensity ratios of different transitions and for different elements (Figs. \\ref{fig:rat} and \\ref{fig:rat_m}), and also of derived physical parameters, including the ionization fraction (Figs. \\ref{fig:par} and \\ref{fig:par_m}).  An accurate treatment of the radiative losses allows us to more realistically compute the full ionization state of the plasma, and thus accurate line emission for each ion, in the approximation of a five-level atom. This, however, requires significant additional computing power that is not always available. An alternative is the use of the simplified cooling model, following only the ionization state of hydrogen, that proved to give good indications on the perturbation amplitude and transverse magnetic field of our model. This simpler but much faster strategy may be used for a preliminary exploration of the parameter space, being followed by simulations with detailed cooling only in the areas of interest. Both cooling models were discussed and compared in Te\\c{s}ileanu et al. (\\cite{TA08}), and in the present work we extended the comparison to line ratio diagnostics and application to a real HH object.  Grid resolution proved to be a critical parameter for the line ratios estimations from MHD simulations. The physical dimensions of the computational zones have to not exceed about $10^{11}$ cm in order to accurately capture the evolution of the physical parameters and emission properties in the post-shock zones. This is a crucial aspect for future 2D simulations that will have to fulfill these requirements. This result can be generalized for the typical conditions encountered in shocks traveling through protostellar jets, as the variability in these conditions could not change the numbers by an order of magnitude. Weaker shocks (shocks evolving from smaller initial perturbations) or lighter shocks (lower initial density) will be less stringent on the resolution requirements, while stronger and/or denser shocks will add a factor to the present results.  We can summarize the general picture as follows: i) when one is interested in the distribution in space of the jet physical parameters such as density, temperature, etc., the details of the assumed cooling function matter very little (Te\\c{s}ileanu et al. \\cite{TA08}) provided the numerical resolution suffices to minimize numerical dissipation effects; ii) when one considers the detailed shock structure, numerical resolution is more important, but less so for the behavior of the integrated emission line luminosities (\\cite{Raga07}); iii) when one discusses the line intensity ratios, a finer resolution must be achieved, while the details of the cooling function have effects that differ according to the line considered, in the present case from about 10\\% up to a factor of two."
    },
    "0910/0910.4894_arXiv.txt": {
        "abstract": "{} {We aim to provide constraints on evolutionary scenarios in clusters. One of our main goals is to understand whether, as claimed by some, the cool core/non-cool core division is established once and for all during the early history of a cluster. } {We employ a sample of $\\simeq$ 60 objects to classify clusters according to different properties: we characterize cluster cores in terms of their thermo-dynamic and chemical properties and clusters as a whole in terms of their dynamical properties.} {We find that: I) the vast  majority of merging systems feature high entropy cores (HEC); II) objects with lower entropy cores feature  more pronounced metallicity peaks than objects with higher entropy cores. We identify a small number of medium (MEC) and high (HEC)  entropy core systems which, unlike most other such objects, feature a large central metallicity.  The majority of these outliers are mergers, i.e. systems far from their equilibrium configuration. } { We surmise that medium (MEC) and high (HEC) entropy core systems with a large central metallicity recently evolved from low entropy core (LEC) clusters that have experienced a heating event associated to AGN or merger activity.} ",
        "introduction": "\\label{sec: intro} The classification of objects is an important step in the construction of viable physical models. This is particularly true for disciplines like astrophysics where, due to the impossibility of performing measurements under  a set of rigorously controlled conditions, evolutionary schemes are inferred primarily by comparing  properties observed in different objects. Galaxy clusters are no exception to this rule, as for other astrophysical sources, much of the early work has concentrated on establishing a taxonomical framework. In the optical band classification schemes are based on the richness \\citep{abell58,zwicky68} and on the morphological properties which have been found to correlate with the dynamical state of the systems \\citep{abell65,abell75a}. In X-rays, classification attempts are generally focused on core properties for the rather obvious reason that cores are the regions more easily accessible to observations. Most workers concentrate on defining indicators discriminating between cool core (hereafter CC) and non-cool core (hereafter NCC) systems; these indicators are typically based on estimates of the intensity of the surface brightness peak  \\citep{Vikh07}, the temperature \\citep[e.g.][]{Sn06}, the cooling time \\citep[e.g.][]{baldi07} or the entropy \\citep[e.g.][]{Ca09} of the central regions of clusters. There have been attempts to derive dynamical properties of cluster ensembles from X-rays, the best known example is perhaps that of the power-ratios technique \\citep{buote95,buote96}. Comparison of  core with morphological properties,  as identified with the power-ratios technique, shows that more disturbed systems tend to have less defined cores \\citep{buote96,bauer05}. One should however keep in mind that the characterization of the degree of relaxation of clusters through X-ray morphology is limited by projection effects \\citep[e.g.][]{jeltema08} and that the limited amount of information available at large cluster radii, beyond $\\simeq 0.2 R_{180}$, provides a further complication. There have been some attempts to compare dynamical and core properties on individual objects \\citep[e.g.~A2034,][]{Ke03}, sometimes using information from different energy bands \\citep[e.g.~A1644,][]{Re04}.  Classification schemes based on chemical properties of clusters have enjoyed considerably less attention. We have known for some time now that CC clusters, formerly known as cooling-flow clusters, feature significant central abundance excesses \\citep{grandi01} and that the amount of iron associated with these excesses is consistent with being produced from the BCG galaxy invariably found at the center of CC systems \\citep{grandi04}. Interestingly, there has been no attempt so far to classify clusters by simultaneously making use of thermo-dynamical and chemical properties. In this paper we employ a medium size sample, $\\simeq$~60 objects, to address the issue of cluster classification from various angles, more specifically we will: 1) define an entropy indicator to classify cluster cores with respect to outer regions; 2) provide a dynamical classification based on radio, optical and X-ray properties; 3) compare core and dynamical properties; 4) compare, for the first time, the entropy based classification with a chemical classification. As we shall see, our classification work will allow us to gain considerable insight in how cluster dynamical, thermo-dynamical and chemical properties relate to each other. Another interesting result that will emerge from our analysis is the relevance of outliers, i.e. objects that fall outside of the distributions defined by the majority of systems, in constraining the evolutionary processes that shape galaxy clusters.  The breakdown of the paper is the following. In Sects.~\\ref{sec: sample} and \\ref{sec: analysis} we describe respectively the sample selection and the data analysis. In Sect.~\\ref{sec: entropy} we provide an account of how our entropy and cooling time indicators have been constructed. In Sect.~\\ref{sec: class sch} we describe two classification schemes based respectively on core and dynamical properties while in Sect.~\\ref{sec: entr vs alt} we compare them with the  entropy based classification scheme defined in Sect.~\\ref{sec: entropy}. In Sect.~\\ref{sec: Z scheme} we define a classification scheme based on chemical properties and compare it to the other classification schemes discussed in the paper. Finally in Sect.~\\ref{sec: concl} we summarize our main findings.   Quoted confidence intervals are 68\\% for one interesting parameter (i.e. $\\Delta\\chi^2$~=~1), unless otherwise stated. All results assume a $\\Lambda$CDM cosmology with $\\Omega_\\mathrm{m} = 0.3$, $\\Omega_\\Lambda = 0.7$, and $H_0$~=~70~km~s$^{-1}$~Mpc$^{-1}$. ",
        "conclusions": "\\label{sec: concl}  The main results presented in this paper may be summarized as follows. \\begin{itemize}  \\item{} We have constructed an indicator of the entropy of the core relative to that of the cluster, the pseudo entropy ratio $\\sigma$. Our indicator is robust, in the sense that somewhat different choices of the quantities from which the ratio is computed result in very similar values of $\\sigma$. The indicator is also relatively parsimonious, in the sense that it may be constructed from data of moderate statistical quality.  \\item{} The classification of clusters based on the entropy indicator improves upon the traditional classification scheme based on the presence or absence of a temperature drop in the core. Conversely classification schemes based on the central cooling time appear to be essentially equivalent to ours.  \\item{} A comparison between the entropy based classification scheme and a classification scheme based on dynamical properties shows that the large majority of merging systems are characterized by large entropy ratios. Only 2 of our merging systems feature a low entropy core (LEC) and in both cases we were able to establish, with reasonable certainty, that the effects of the merger have not reached the core. Our findings are at variance with those presented by \\citet{Mc04} who do find evidence of non relaxed cooling flow systems and relaxed non-cooling flow systems. We have compared our entropy and dynamical classification schemes with those in \\citet{Mc04} for the 15 objects that are common to both samples finding that the 3 cases where our classifications do not agree likely result from misclassifications by \\citet{Mc04}.   \\item{} We find that mean abundance profiles for our 3 entropy classes, namely low entropy core (LEC), medium entropy core MEC, and high entropy core (HEC), may be divided in 3 regions. In the outer region, between 0.2 and 0.4~$R_{180}$, all 3 profiles are consistent with being flat and with one another. In the core region, within 0.1~$R_{180}$, all classes feature an excess with respect to the mean value found in the outer regions. The excess is strongest for LEC, somewhat weaker for MEC, and weakest for HEC clusters. Between 0.1 and 0.2~$R_{180}$ the profiles for the three classes are roughly consistent with one another and, at least for LEC and MEC systems, show a significant excess with respect to the mean value measured in the outskirts.   \\item{} We find that objects with stronger pseudo-entropy gradients have  more pronounced metallicity peaks. This results tells that, baring  a few exceptions, the gas that is more enriched in metals  also happens to be the one featuring the lowest entropy.   \\item{} We have identified a small number of medium and high  entropy core systems with a large central metallicity.  The majority of these objects have been classified as mergers, i.e. as systems far from their equilibrium configuration.  We surmise that these systems evolved from low entropy core clusters that have experienced a heating event. We have examined  simulation based claims that conflict with our conjecture finding they are flawed both on the observational and the theoretical side.   \\end{itemize}  In an upcoming paper (Rossetti \\& Molendi 2009) we will investigate further the issue of medium and high  entropy core systems with a large central abundance; we will do this by performing bi-dimensional analysis of a smaller sample of bright and nearby clusters."
    },
    "0910/0910.4546_arXiv.txt": {
        "abstract": "Many multi-planet systems have been discovered in recent years. Some of them are in {mean-motion} resonances (MMR). Planet formation theory was successful in explaining the formation of 2:1, 3:1 and other low resonances as a result of convergent migration. However, higher order resonances require high initial orbital eccentricities in order to be formed by this process and these are in general unexpected in a dissipative disk. We present a way of generating large initial eccentricities using additional planets. This procedure allows us to form high order MMRs and predict new planets using a genetic {$N$}-body code. ",
        "introduction": "Of the recently discovered 322 extrasolar planets, at least 75 are in multi-planet systems (Schneider \\cite{schneider}). About 10\\% of these are confirmed to be in or very close to a resonant configuration.  Two planets can easily be locked into a low order resonance such as 2:1 or 3:2 by dissipative forces. This process is well understood (see eg. {Lee \\& Peale \\cite{leepeale}}, Papaloizou \\& Szuszkiewicz \\cite{papaloizou}). However, period ratios of multi-planetary systems seem to be non-uniformly distributed up to large period ratios of order 5.  Furthermore, there is an observational bias against the detection of resonant systems in general (Anglada-Escude \\etal{}  \\cite{anglada}). One should note that a period ratio of e.g. approximately 5 does not imply that the system is indeed in a 5:1 {mean-motion} resonance. We need long term, high precision measurements to determine exactly in which dynamical state the planets are. For most observed systems, this has not been achieved yet. However, we hope that the increasing number of detected planets will enable us to draw a conclusion in a statistical manner in the near future.   These high order resonance candidates cannot be explained by adiabatic migration only. The capture requires a high initial eccentricity which is not expected to be present in the system shortly after formation of the proto-planet cores.   In {Figure} 1, the final period ratio of two migrating Jupiter mass planets is plotted as a function of the initial eccentricity of the outer planet and the eccentricity damping parameter $K$ which is defined as the ratio of the timescales $\\tau_a$ and $\\tau_e${, being the migration timescale and the eccentricity damping timescale, respectively.} The color encodes the final period ratio at the end of the simulation. At the beginning the planets are put far away from each other. The outer planet migrates inwards on a timescale of $\\tau_a=25000$~years. One can see that a small initial eccentricity always results in a 2:1 mean motion resonance. At slightly higher eccentricities ($e \\sim 0.05$) the favored outcome is a 3:1 resonance. At large eccentricities ($e \\sim 0.15$) it is possible to get high resonances like 4:1 and 5:1. Note that if $K$ is large, the eccentricities get damped before the planets {are captured} into resonance and therefore favor lower order resonances.   \\begin{figure}[htb] \\begin{center} \\includegraphics[angle=270,width=0.7\\columnwidth]{fig0} \\caption{This figure shows the final period ratio of two convergently migrating Jupiter mass planets as a function of the initial eccentricity $e_{init}$ of the outer planet and the ratio of damping timescales $K$ (see text). {Red, as indicated on the color bar, corresponds to a 2:1, green to 3:1, blue to 4:1 and yellow to 5:1 commensurabilities respectively. Black indicates that the period ratio is not close to an integer value or the planets scattered within a short period of time.}} \\end{center} \\end{figure}   \\newpage ",
        "conclusions": "It is possible to model the observed orbital parameters of systems with high order commensurabilities in a natural way by postulating additional planets. The {solution presented here} does not rely on the random effects of {planet-planet} scattering. The prediction of new planets can result in a test for planet formation theories and provides attractive targets for observers.  An interesting point that hasn't been discussed here is the question whether this scenario is stable to small perturbations. These perturbations could result from a turbulent disc or other planets (see e.g. Rein \\& Papaloizou \\cite{rein}).  Further work is underway, including hydrodynamical simulations confirming the solutions found by the {$N$}-body code.  \\newpage"
    },
    "0910/0910.1363_arXiv.txt": {
        "abstract": "We model the optical to far-infrared SEDs of a sample of six type-1 and six type-2 quasars selected in the mid-infrared. The objects in our sample are matched in mid-IR luminosity and selected based on their {\\it Spitzer} IRAC colors. We obtained new targeted {\\it Spitzer} IRS and MIPS observations and used archival photometry to examine the optical to far-IR SEDs. We investigate whether the observed differences between samples are consistent with orientation-based unification schemes. The type-1 objects show significant emission at 3 \\micron. They do not show strong PAH emission and have less far-IR emission on average when compared to the type-2 objects. The SEDs of the type-2 objects show a wide assortment of silicate features, ranging from weak emission to deep silicate absorption. Some also show strong PAH features. In comparison, silicate is only seen in emission in the type-1 objects. This is consistent with some of the type-2s being reddened by a foreground screen of cooler dust, perhaps in the host galaxy itself. We investigate the AGN contribution to the far-IR emission and find it to be significant. We also estimate the star formation rate for each of the objects by integrating the modeled far-IR flux and compare this with the SFR found from PAH emission. We find the type-2 quasars have a higher average SFR than the type-1 quasars based on both methods, though this could be due to differences in bolometric luminosities of the objects. While we find pronounced differences between the two types of objects, none of them are inconsistent with orientation-based unification shemes. ",
        "introduction": "The study of various types of AGN can provide insight into their origin and evolution. AGN are classified according to their optical emission lines as type-1 or type-2. Type-1 AGN have an \\textquotedblleft unobscured\\textquotedblright~view of the hot accretion disk and broad emission lines originating from fast moving clouds that surround a supermassive black hole. Type-2 objects are known as \\textquotedblleft obscured\\textquotedblright~AGN, because no broad lines or continuum emission from the accretion disk are observed, yet the spectra still show narrow emission lines similar to those of the type-1 objects.  The interpretation of the differences between type-1 and type-2 quasars (the most luminous AGN) remains an open issue. The orientation-based unification scheme proposes the two types of objects are intrinsically the same, but appear different observationally due to the presence of an obscuring torus of dust close to the central power source (see Antonucci 1993). Type-1 objects are seen from an orientation which allows an unobscured view of the central source. This allows the broad line region of the quasar to be observed. However, type-2 objects are seen through the dusty torus material, which hides the central continuum source and broad-line region. Evidence that this model is at least partly correct comes from observations of polarized broad-line emission in galaxies classified as type-2 quasars (\\eg~Zakamska et al. 2005). But this model does not provide insight into the origin of quasar activity.  An alternative model proposes that quasar activity is triggered by galactic mergers and evolves through an obscured stage \\citep{S88}. As the galaxies merge, gas and dust are stirred up, triggering quasar activity and star formation. Initially the quasar would be obscured by the dust, but over time will become unobscured as dust and gas are blown away from the nuclear region by quasar and/or supernovae driven winds and radiation pressure. Indeed, several host galaxy studies have been made which show that luminous AGN are hosted by massive, mainly spheroid-dominated galaxies, with a significant fraction, but not all, showing signs of a relatively recent merger and associated star formation (\\eg~Dunlop et al. 2003; Canalizo et al. 2007; Bennert et al. 2008). Thus the evolution model predicts that the host galaxies of type-2 quasars should have more recent star forming activity than those of type-1s. Although the discovery of polarized broad lines is a powerful argument in favor of the orientation-based model, it does not rule out an evolutionary link.  The obscuring dust in a host galaxy would be apparent in its observed SED. Dust close to the AGN would be heated by the central power source and would thermally re-radiate in the mid-IR (100 - 1000K), dominating the rest-frame SED from $\\sim$2 - 40 \\micron. This high temperature dust emission makes the mid-IR SEDs of both optically-obscured and optically-unobscured AGN quite distinct from those of either normal or starburst galaxies, whose dust temperatures are always $\\leq$ 100K. By modeling the various components of the SEDs, we can compare the relative amounts of ``cool'', \\textquotedblleft warm\\textquotedblright, and \\textquotedblleft hot\\textquotedblright~dust emission present in the host galaxies. This will give us some idea of what mechanism dominates the heating of the dust (star formation or AGN), and how much of a role orientation might play in the observed SEDs.  In this paper, we present modeled fits to the SEDs of a sample of type-1 and type-2 quasars of similar redshift and luminosity. \\citet{Lacy07} presented preliminary fits to the type-2 quasars in this paper. We also measure the objects' far-IR luminosities, PAH luminosities, and estimate their star formation rates. We then compare the results with starburst galaxies. We describe the sample selection in \\S2, and the observations and reductions in \\S3. In \\S4 we present the analysis and results including our modeling procedure and resultant spectra, a color-color diagram, and star formation rate (SFR) calculations. Finally, we summarize our results in \\S5. Throughout this paper we adopt the cosmological parameters of H$_{0}$=71 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{M}$=0.27 and $\\Omega_{\\Lambda}$=0.73. ",
        "conclusions": "We obtained mid-IR spectra and far-IR photometry for six type-1 and six type-2 mid-IR selected quasars. We modeled the SEDs of our sample from the optical to the far-IR using available SDSS, 2MASS and IRAC photometry, in addition to our IRS spectra and MIPS photometry. We compared the modeled SEDs of the quasars and calculated SFRs based on PAH and far-IR emission. Our main results are:   (1) The type-1 quasar mid-IR spectra show a featureless continuum with significant 3 \\micron~emission. The full optical through far-IR SEDs are typically fit well with an optical quasar component, a 1000 K modified blackbody, a mid-IR power-law ($\\sim3 - 30$\\micron), and a cool dust far-IR  modified blackbody (45 K). Three of the objects also show evidence of warm small grain dust fit with a power-law ($\\sim25 - 65$\\micron).   (2) The type-2 quasar mid-IR spectra show a variety of PAH emission and silicate absorption strengths. The full optical through far-IR SEDs are fit well using a stellar population in the optical, a mid-IR power-law ($\\sim3 - 30$\\micron), varying amounts of a PAH emission template and silicate absorption in the mid-IR, warm small grain power-law emission ($\\sim25 - 65$\\micron) as well as a cool dust modified blackbody (45 K).   (3) Our averaged and normalized modeled spectra show that type-2 quasars exhibit more far-IR emission than type-1 quasars. The excess is roughly equivalent to the deficit in the optical with respect to the type-1 quasars, which is consistent with the orientation hypothesis. In this scenario the increased far-IR emission arises from obscuring dust that is reradiating absorbed light.   (4) We measured L$_{PAH}$ by integrating our scaled template for each object. We compared the typical L$_{PAH}$ to L$'_{FIR}$ ratio of our sample to that of starburst galaxies that have had L$_{PAH}$ measured in a similar manner as the method we used. Within our sample, we found that the type-2 objects have a larger ratio than the type-1 objects. We note that the difference between the two is not significant. Considering our overall sample, we found that the ratio is smaller than for the starburst galaxies \\citep{M07}. This result suggests some far-IR emitting dust can be warmed by the central AGN rather than by young stars. However, a more direct comparison between AGN and starburst galaxies using larger samples would be needed to further quantify the amount of AGN-heated far-IR emitting dust.   (5) We calculated the SFRs of our sample using two methods. First we integrated the warm small grain and cool modified blackbody emission (excluding all other components) of the modeled SEDs and used the \\citet{K98} relation to calculate the SFRs. We also converted the L$_{PAH}$ to L$_{FIR}$ \\citep{M07} to measure the SFRs based on the PAH emission only. In each case the type-2 objects had higher SFR on average than the type-1 objects. However, the difference is not significant given the scatter.  (6) We integrated our modeled SEDs from $8 - 1000$~\\micron~and found that seven of the 12 can be considered ULIRGs (log(L$_{IR}/$L$_{\\odot}$) $\\geq 12$), while four others have log(L$_{IR}$/L$_{\\odot}$) $\\geq 11.75$. The sole exception is SST 1721+6012, which has log(L$_{IR}/$L$_{\\odot}$) $= 11.48$. While our sample appears to have IR luminosities similar to ULIRGs, the quasars exhibit SFRs more typical of LIRGs. We found the AGN can heat most of the dust in the mid-IR and even some in the far-IR. Thus measuring SFR in quasars from the full $8 - 1000$~\\micron~SED overestimates the amount of dust warmed by young stars and the true SFR.   (7) The modeling results suggest that type-1 quasars can have mid-IR fluxes that have significant contributions from the optical power-law and 3 \\micron~bump.  This 3 \\micron~component is not required for the type-2 SEDs, and the hot dust most likely originates near the sublimation radius of the AGN. If this component is present in the type-2 quasars, it must be at least partially obscured by cooler dust, either a dusty torus or dust in the host galaxy. \\citet{Lacy07} showed that the host galaxies of the type-2 quasars have a range of inclinations.  Therefore, the obscuration of the 3 \\micron~bump, if it is intrinsically present, should be from material close to the nucleus and independent of galaxy orientation.  While our analysis does not distinguish definitively between orientation and evolution based unification schemes, it does directly address the dust content, ongoing star formation in the host galaxies and the AGN contribution to the far-IR emission. Our results are consistent with the orientation hypothesis, but we cannot rule out an evolutionary connection. The star formation based on PAH measurements is a larger contributor to the far-IR luminosity than the AGN for type-2 objects compared to type-1 objects. This difference is only suggestive due to the scatter of both L$_{PAH}$ and L$'_{FIR}$ in our sample and could arise from intrinsic differences in the bolometric luminosities of the objects. The L$_{PAH}$/L$'_{FIR}$ ratio in both type-1 and type-2 quasars is lower than that of starburst galaxies, suggesting that some emission in the far-IR is from dust heated by the AGN and not star formation. This is an area for continued research, as our analysis was based on few objects."
    },
    "0910/0910.2395.txt": {
        "abstract": "We %attempt to quantify the rapid variations in X-ray brightness (``flares'') from the extremely massive colliding wind binary \\ec\\ %as seen during the past three orbital cycles by \\rxte. The observed flares tend to be shorter in duration and more frequent as periastron is approached, although the largest ones tend to be roughly constant in strength at all phases.  Plausible scenarios include (1) the largest of multi-scale stochastic wind clumps from the LBV component entering  %or pushing on and compressing the hard X-ray emitting wind-wind collision (WWC) zone, (2) large-scale corotating interacting regions in the LBV wind sweeping across the WWC zone, or (3) instabilities intrinsic to the WWC zone. The first one appears to be most consistent with the observations, requiring homologously expanding clumps as they propagate outward in the LBV wind and a turbulence-like power-law distribution of clumps, decreasing in number towards larger sizes, as seen in Wolf-Rayet winds. ",
        "introduction": "The extremely massive binary star \\ec\\ consists of a Luminous Blue Variable (LBV) primary (star A) coupled with a hot, fast-wind secondary (star B) which is probably an evolved O star \\citep{2005ApJ...624..973V} or a Wolf-Rayet star \\citep{2002AA...383..636P}. Orbital and stellar properties gleaned from the literature are given in Table \\ref{tab:params}. These values are considered to be the best current values, although overall relatively uncertain. Star A has gone through at least one giant LBV eruption, which occurred in the 1840s \\citep{1997ARA&A..35....1D}.   \\begin{deluxetable}{lcl} \\tablecaption{Adopted Parameters for the \\ec\\ Binary System} \\tablehead{ \\colhead{Parameter} & \\colhead{Value} & \\colhead{Reference} } \\startdata $P$(orbit) & $2024.0 \\pm 2$ d = $5.541 \\pm 0.006$ yr & \\cite{2005AJ....129.2018C}\\\\ T$_\\circ$\\tablenotemark{a} & HJD $2450800\\pm3$ d (=1997.9618 UT) & \\cite{2005AJ....129.2018C}\\\\ $i$ & $45^{\\circ}$ &  \\cite{2008MNRAS.388L..39O}\\\\ $e$ & 0.9 &  \\cite{2008MNRAS.388L..39O}\\\\ $\\omega$ & 243$^{\\circ}$  &  \\cite{2008MNRAS.388L..39O}\\\\ $a$ & 15.4 AU & \\cite{2001ApJ...547.1034C}\\\\ $D$(periastron) & 1.54 AU & \\\\ $D$(apastron) & 29.3 AU & \\\\ $M_{A}$ & 90  M$_{\\odot}$ & \\cite{2001ApJ...553..837H}\\\\ $M_{B}$ & 30 M$_{\\odot}$ & \\cite{2005ApJ...624..973V}\\\\ ${\\dot M}_{A}$ & 10$^{-3}$ M$_{\\odot}$ yr$^{-1}$ & \\cite{2001ApJ...553..837H}\\\\ ${\\dot M}_{B}$ & 10$^{-5}$ M$_{\\odot}$ yr$^{-1}$ & \\cite{2002AA...383..636P}\\\\ $V_{\\infty,A}$ & 500 km/s & \\cite{2001ApJ...553..837H}\\\\ $V_{\\infty,B}$& 3000 km/s &  \\cite{2002AA...383..636P}\\\\ $R_{A}$ & 0.28 AU & \\cite{2001ApJ...553..837H}\\\\ %$v_{breakup}$(A) & 300 km/s \\\\ \\enddata \\tablenotetext{a}{T$_{\\circ}$ is the Heliocentric Julian Day Number at the start of the X-ray minimum in 1997-1998, which is likely close to the time of periastron passage.} \\label{tab:params} \\end{deluxetable}% %P(orb) = 2024.0 +/- 2 d = 5.54 +/- ?? yr %i = 41 +/-??$^{\\circ}$ %e = 0.9 +0.05/-0.1 %$\\omega$ = 270 +0/-90$^{\\circ}$ %T$_o$ = HJD 2 4??.  +/-?? %a = 15.4 +/- ?? AU %r(periastron) = 1.54 AU %r(apastron) = 29.3 AU %M(A) = 90 +/- ?? M$_{\\odot}$ (Hillier\u0085) %M(B) = 30 +/- ?? M$_{\\odot}$ (Verner et al. \u0085) %$v_\\infty$(A) = 500 km/s %$v_\\infty$(B) = 3000 km/s %${\\dot M$(A) = 10$^{-3}$ M$_{\\odot}$ yr$^{-1}$ %${\\dot}M$(B) = 10$^{-5}$ M$_{\\odot}$ yr$^{-1}$  \\ec\\ generates X-ray emission in the $2-10$ keV range due to the collision of star A's dense, slow wind with the thinner, faster wind of star B \\citep{1999ApJ...524..983I}.  On timescales of years, the X-ray emission varies in a manner similar to that expected from a long-period, highly eccentric, colliding-wind binary, with $L_X \\propto 1/D$ (where $D$ = orbital separation), for adiabatic conditions as is most likely the case in such a long-period system \\citep{usov92}, though near periastron passage radiative effects \\citep{2009MNRAS.tmp..279P} may become important. This trend is broken by a broad atmospheric eclipse in the X-ray light-curve, which %occurs begins near periastron passage when the dense inner wind of Star A starts to block out most of the X-ray emission arising in the bow-shock head region of the wind-wind collision (WWC) zone. Along with this long-term variation, short term variations in \\ec's X-ray brightness, or ``flares'' \\citep{1997Natur.390..587C,1998NewA....3..241D,1999ApJ...524..983I}, have been observed by \\rxte\\ for more than two full orbital cycles from 1996.2 up to at least 2009.0, near the date of this writing.  These flares are not expected in standard colliding wind models, and their origin is currently unknown.  We examine here the characteristics of these flares and discuss possible physical mechanisms for their production. ",
        "conclusions": "Basically, all three models have varying degrees of problems.  In addition, there is the shared problem that the orbital parameters of \\ec\\ are not very well known. But despite the inherent subjective problem of actually defining the flares in \\ec\\, we believe that the clump model is the simplest (and therefore most likely) explanation yet found for the flares.  LBV stellar pulsations on a time scale of $\\sim$85 d seem less likely because of the lack of a true oscillation signature. % [TO TONY: I DON'T! - I THINK THERE ARE SCALE PROBLEMS WITH THE CLUMP MODEL, IN ADDITION TO THE PHASE DEPENDENCE OF THE FLARE SEPARATION WHICH IS STILL NOT ADEQUATELY ADDRESSED!].   Clumping on relatively large scales has been invoked to explain the X-ray lines in O stars \\citep{2003A&A...403..217F}. \\cite{2007A&A...476..335W} use similar model parmeters to explain flaring in X-ray binaries. Thus, wind collisions have some similarity to massive X-ray binaries, in which relatively hard X-ray flares are seen as a likely result of clumpy winds being accreted by the compact component.  If so, then one should should see X-ray flares in other wind-wind colliding systems, such as the prototype WR + O system WR140 = HD 193793, WC7pd + O5, P = 7.94 yr, e = 0.88.  The X-ray lightcurve of WR140 is also being monitored by \\rxte. If it shows flares (in fact there is some evidence for absorption dips possibly from clumps), then they are far less obvious than those in \\ec.  One difference between these two systems, though, is that the two colliding winds in WR140 are both very fast and may not contain the %very large prominent clumps that LBV winds appear to have \\citep{2007A&A...469.1045D}, %(Davies et al. 2007), leading to the giant X-ray flares that make \\ec\\ so spectacular.   %In this way, wind collisions have some similarity to massive X-ray binaries, in which relatively hard X-ray flares are seen as a likely result of clumpy winds being accreted by the compact component %\\citep[e.g.][]{2007A&A...476..335W}.  If so, then one should see X-ray flares in other wind-wind colliding systems, such as the prototype WR + O system WR140 = HD 193793, WC7pd + O5, P = 7.94 yr, e = 0.88.  The X-ray lightcurve of WR140 is also being monitored by \\rxte. If it shows flares (in fact there is some evidence for absorption dips possibly from clumps), then they are far less obvious than those in \\ec.  One difference between these two systems, though, is that the two colliding winds in WR140 are both very fast and may not contain the %very large %prominent clumps that LBV winds appear to have \\citep{2007A&A...469.1045D}, %%(Davies et al. 2007), %leading to the giant X-ray flares that make \\ec\\ so spectacular.    %\\subsection{Strong flare interval}  %\\subsection{Flare times} %In order to test for whether peak flare times are correlated, we took the flare amplitudes and times of peak emission as given in \\citet{2005AJ....129.2018C}."
    },
    "0910/0910.0773_arXiv.txt": {
        "abstract": "We have investigated spectral variation of the Seyfert 1 galaxy  MCG-6-30-15 observed with Suzaku in  January 2006 for three separate  periods spreading over fourteen days.  We found that the time-averaged continuum energy spectrum between 1 keV and 40 keV  can be approximated with a spectral model composed of the direct power-law component, its reflection component, two  warm absorbers with different ionization states, and neutral absorption. We have taken two approaches to study its spectral variation at various timescales:  The first approach is to make intensity-sliced spectra and study correlation between the intensity and spectral shape. The second approach is to study spectral changes between the intervals when the source flux is above (``bright state'') and below (``faint state'') the average for  fixed  time-intervals. In both approaches, we found a clear correlation between the intensity in the 6 -- 10 keV band and the spectral ratio of 0.5 -- 3.0 keV/6.0-- 10 keV. Such a spectral variation requires change of the apparent  slope of the  direct component, whereas the shape and intensity of the reflection component being invariable. The observed apparent spectral change is  explained by variation of the ionization degree  of one of the two warm absorbers due to intrinsic source luminosity  variation.  Current results suggest that the warm absorber has a critical role to explain the observed continuum spectral shape and variation of MCG-6-30-15, which is essential to constrain parameters of the putatively  broadened iron line emission feature. ",
        "introduction": "The Seyfert 1 galaxy MCG-6-30-15 has been extensively studied with   ASCA (Iwasawa et al. 1996,1999; Matsumoto et al. 2003), BeppoSAX (Guainazzi et al. 1999), Rossi X-ray Timing Explorer (Lee et al. 1999, Vaughan $\\&$ Edelson 2001), Chandra (Lee et al. 2001, Young et al. 2005, Gibson et al. 2007), XMM-Newton (Fabian et al. 2002, Vaughan $\\&$ Fabian 2004, Ponti et al. 2004), Suzaku (Miniutti et al. 2007), and a combination of Suzaku, Chandra and XMM-Newton (Miller, Turner, and Reeves 2008). ASCA discovered a broad and skewed emission line feature around 5--7 keV in the spectrum of MCG-6-30-15  for the first time   (Tanaka et al. 1995). The iron line is emitted as a part of the reflection spectrum from the accretion disk,  irradiated by a primary continuum  of  the central engine. Fabian et al. (1989) suggested  that the iron line due to  disk reflection is in general broadened and skewed by the Doppler effects and gravitational redshift if the line is emitted in the vicinity of the black hole. In fact, the iron line profile observed with ASCA is well explained with such a   ``disk-line'' model (Tanaka et al.\\ 1995).  Meanwhile, Fabian et al. (2002) and Matsumoto et al. (2003) reported significant {\\em invariability}\\/ in the  iron line energy band of MCG 6-30-15 in XMM-Newton and ASCA observations, respectively. Inoue and Matsumoto (2003) pointed  out that the absorbed spectrum due to photoionized warm absorbers mimics the shape of the strongly red-shifted iron line, and that variation of the warm absorbers may explain such invariability in the iron line region. On the other hand, Miniutti and Fabian (2004) proposed that the observed invariability of the iron line may be  explained if the disk reflection and line photon production takes place very close to the black hole, such that the apparent invariability of the iron line is due to general relativistic light-bending effects. In fact, Miniutti et al. (2007) claim that the strong reflection, the broad iron line and  invariability of the iron line observed with Suzaku may be explained with the light-bending model, and suggest that the innermost disk radius extends down to about as low as two gravitational radii. Meanwhile, Nied\\'{z}wiecki \\& \\.{Z}ycki (2008) independently re-examined the light-bending model and concluded that it is indeed possible that a reflected component is very weakly variable compared to the primary emission,  if a source is moving {\\em radially}\\/ very close to the rotating black hole, while whether such a particular movement is plausible or not is another question.   On the other hand, a similarly weakly variable, {\\em apparently broad}\\/ Fe line in NGC 4151 is most likely to  be an artifact of incorrect modelling of the absorbers, and the spectrum is in fact dominated by a narrow emission line (Ogle et al.\\ 2000; de Rosa et al.\\ 2007). As a matter of  fact, Takahashi, Inoue, and Dotani (2002) confirmed, from a very long RXTE monitoring data,  an unambiguous evidence that the reverberation of  the iron line emission takes place   not in the vicinity of the central black hole but very far from the black hole ($\\sim10^{17}$ cm). Therefore, at least in NGC 4151, most iron line photons are produced very far from the black hole and thus invariable on a timescale of $\\sim10^5$ sec.  In MCG-6-30-15 too,  Young et al. (2005) reported a weak and narrow emission  line at 6.4 keV, which indicates that some reflection does take place in a distant material, e.g., the narrow-line region or the pc-scale torus. However, equivalent width of the narrow line in MCG-6-30-15 ($\\sim18$ eV) is much smaller than in NGC 4151 ($\\lesssim$ 2 keV; Takahashi, Inoue and Dotani 2002). Therefore, invariability of the apparently broadened iron line feature in MCG-6-30-15 has yet to be explained.  The question in MCG-6-30-15 is thus whether the Fe line is broad and its lack of variability is a result of the relativistic effect, or it is rather narrow, comes from a distant and extended medium, which would naturally smear the  variability.  The crucial element to address the problem is to model the underlying continuum spectrum and study its variation. The broad band continuum in MCG-6-30-15 is complex, with strong modification from warm absorbers (Reynolds et al.\\ 1995, Otani et al.\\ 1996, Lee et al.\\ 2001, Sako et al.\\ 2003, Turner et al.\\ 2003, Turner et al.\\ 2004, Young et al.\\ 2005, Miller et al.\\ 2008). Moreover, the continuum spectral shape is variable, and the spectral variation is suggested to be  driven mostly by variations of the complex warm absorbers (Miller et al.\\ 2008).  In this paper we attempt to comprehensively characterize the spectral variability of the source in a model independent manner as much as possible. We characterize  spectral variations at different timescales, as well as at different source flux levels. We use Suzaku data taken in Jan 2006 (total exposure time is 339 ksec) for this purpose. Suzaku's broad energy coverage (0.6--40 keV),  excellent spectral resolution and a large effective area, make it an  ideal instrument to study spectral variation of moderately bright AGNs such as MCG-6-30-15. Our goal is to find a reasonable spectral model of MCG-6-30-15 which is able to explain the observed  spectral variation at  various timescales with minimum  numbers of parameters, and to study effects of warm absorbers. Such a reasonable spectral model is very needed to disentangle the tough disk-line problem, since the broad iron line feature is affected by the choice of continuum spectral models. ",
        "conclusions": "\\begin{figure} \\begin{center} \\FigureFile(80mm,80mm){figure10a.ps} \\FigureFile(80mm,80mm){figure10b.ps} \\end{center} \\caption{Variation of the power-law normalization and depth of the O{\\footnotesize VIII} edge to describe the spectral variation in  0.6 -- 1.6 keV, as functions of the timescales.} \\label{below1keVvariation} \\end{figure}   \\begin{figure}[htbp] \\begin{center} \\FigureFile(80mm,80mm){figure7.ps} \\end{center} \\caption{Unfolded energy spectra for the eight intensity-sliced spectra calculated from the model only varying the power-law normalization and ionization degree of the lower-ionized warm absorber. } \\label{data12} \\end{figure}   \\begin{figure}[htbp] \\begin{center} \\FigureFile(80mm,80mm){figure8a.ps} \\FigureFile(80mm,80mm){figure8b.ps} \\FigureFile(80mm,80mm){figure8c.ps} \\end{center} \\caption{Top: Relation between the variation time scale and ionization degree for the bright and faint spectra. Middle: Relation between the variation time scale and power-law normalization for the bright and faint spectra. Bottom: Relation between the ionization degree and power-law normalization for the bright, faint and sliced spectra. } \\label{xi_variation} \\end{figure}  \\begin{figure}[htbp] \\begin{center} \\FigureFile(80mm,80mm){figure11.ps} \\end{center} \\caption{Ratio of our best-fit model (black) and the observed data (red) to a power-law function.  The power-law slope is determined using the data only in 3.0--4.0 keV and 7.5--12 keV.  Note that our model includes a moderately broad iron emission line at 6.42 keV ($1 \\sigma=$290 eV), but not an extremely distorted one with relativistic effects.} \\label{fig:warm} \\end{figure}   We studied spectral variation of MCG-6-30-15 observed with Suzaku in 2006 January, and found a clear correlation between  the 6 -- 10 keV flux and the spectral  ratio of 0.5 -- 3.0 keV/6.0 -- 10 keV (Figure \\ref{fig:intensityhardness}).  Amplitude of the flux variation and the accompanying spectral change increases when the variation timescales become longer from 5 ksec to 200 ksec  (Figure \\ref{data345}). We fitted the 1 -- 40 keV energy spectra with the spectral model including a cut-off power-law,  two warm absorbers and neutral reflection (Figure \\ref{fig1}). The observed spectral variation was almost equally explained by variation of either photon index of the direct component,  ionization degree of the lower-ionized warm absorber, or column density of the warm absorber, whereas the neutral reflection component was invariable throughout the observation. We could not clearly judge if the intrinsic photon index or  warm absorber parameters are variable, but it is most physically  reasonable to assume that the photoionization degree increases with increasing intensities. We also fit the eight intensity-sliced spectra with the model varying the power-law normalization and the ionization parameter (Table \\ref{table1_slice}). In  Fig.$\\ref{data12}$, we show the incident spectral changes for the intensity sliced spectra calculated from the best-fit parameters.   \\setlength{\\tabcolsep}{2.4pt} \\begin{table*} \\begin{center} \\rotatebox{90}{ \\begin{minipage}[t]{\\textwidth} \\scriptsize \\caption{Results of spectral fitting in 1--40 keV for the intensity sliced spectra when only the power-law  normalization and the ionization degree of the lower-ionized absorber are varied. \\footnotemark[$*$] } \\begin{tabular}{ccccccccc} \\hline state  & slice 1 & slice 2 & slice 3 & slice 4 & slice 5 & slice 6 & slice 7 & slice 8  \\\\ \\hline $N_H$ (${10}^{21}$ ${cm}^{-2}$)  & \\multicolumn{8}{c}{2.3$\\pm$0.1} \\\\ \\hline $N_H$  (${10}^{22}$ ${cm}^{-2}$)  & \\multicolumn{8}{c}{3.3$\\pm$0.9} \\\\ log $\\xi$  & \\multicolumn{8}{c}{3.27$^{+0.04}_{-0.06}$} \\\\ \\hline $N_H$  (${10}^{21}$ ${cm}^{-2}$)  & \\multicolumn{8}{c}{4.1$\\pm$0.3} \\\\ log $\\xi$  &  1.24$^{+0.12}_{-0.10}$  & 1.57$\\pm$0.06 &  1.69$^{+0.06}_{-0.03}$ & 1.70$\\pm$0.06  &  1.74$^{+0.05}_{-0.06}$ & 1.86$^{+0.05}_{-0.06}$ & 1.75$\\pm$0.06  & 2.05$\\pm$0.12  \\\\ \\hline Line E (keV) & \\multicolumn{8}{c}{6.35 (fix)} \\\\ sigma (keV) &  \\multicolumn{8}{c}{0.01 (fix)} \\\\ norm (${10}^{-5}$) & \\multicolumn{8}{c}{1.7$\\pm$0.2} \\\\ EW (eV) & 68$\\pm$6 & 44$\\pm$4  &  40$\\pm$4  &  36$\\pm$3  &  36$\\pm$3  & 35$\\pm$3\t &   33$\\pm$3   & 21$\\pm$2 \\\\ \\hline cutoffpl K (10$^{-2}$) &  0.75$\\pm$0.02 &  1.35$\\pm$0.03  & 1.54$\\pm$0.04  & 1.72$\\pm$0.04  & 1.75$\\pm$0.04\t &  1.77$\\pm$0.04 & 1.89$\\pm$0.05    &  3.23$\\pm$0.12  \\\\ photon index & \\multicolumn{8}{c}{2.04$\\pm$0.02} \\\\ \\hline $E_{cut}$ (keV) & \\multicolumn{8}{c}{160 (fix)} \\\\ cosIncl & \\multicolumn{8}{c}{0.866 (fix)} \\\\ pexrav K (10$^{-2}$)  & \\multicolumn{8}{c}{2.8$^{+0.5}_{-0.3}$} \\\\ \\hline Line E (keV) &  \\multicolumn{8}{c}{2.35$\\pm$0.01} \\\\ sigma (keV) & \\multicolumn{8}{c}{0.01 (fix)} \\\\ norm (${10}^{-5}$) & \\multicolumn{8}{c}{-2.4$\\pm$0.3} \\\\ \\hline reduced chisq (d.o.f) & \\multicolumn{8}{c}{1.24 (1143)} \\\\ \\hline \\multicolumn{8}{@{}l@{}} {\\hbox to 0pt {\\parbox{140mm}{\\footnotesize \\footnotemark[$*$] See the footnote  of Table \\protect\\ref{table0}. }\\hss}} \\label{table1_slice} \\end{tabular} \\end{minipage}} \\end{center} \\end{table*}   In the top panel of Fig.$\\ref{xi_variation}$, we show relation between the variation timescales and ionization degrees of the low-ionized warm absorber. Also, in the middle panel of Fig.$\\ref{xi_variation}$, we show the relations between the variation time scales and the power-law normalizations. It is obvious that both the ionization degree and the power-law normalization for the bright state increases with the variation time-scale, while those for the faint state decreases  with the variation time-scale. In the bottom panel of Fig. $\\ref{xi_variation}$, we show relation between the power-law normalization and the ionization degree of the lower-ionized warm absorber, which indicates a clear correlation as $\\log \\xi $= (1.7$\\pm$0.2) $\\times$ log K + (4.7$\\pm$0.3).  Spectral variation in 0.6 -- 1.6 keV can be also explained by change of the O{\\footnotesize VIII} edge depth due to photoionization (Figure \\ref{below1keVvariation}). Therefore, we suggest that spectral variation of MCG-6-30-15 in 0.6 keV to 40 keV on timescales of 5 ksec to 200 ksec is primarily explained by change of the ionization degree of  warm absorbers due to photoionization, while power-law photon-index and the disk reflection component are invariable.  Note that our model does not include an extremely  broad iron emission line. If the intrinsic line width is allowed to be free, the central line energy is 6.42$\\pm0.06$keV, and the intrinsic line width is 1$\\sigma = 290\\pm60$ eV (Figure \\ref{fig1-2}).  With the combination of the cold reflection, mildly broad emission line and warm absorber, our model can successfully fit the seemingly broad iron line feature (Figure \\ref{fig1-2}). In Figure \\ref{fig:warm}, we show a ratio of our best-fit model to a power-law function, as well as a ratio of the observed spectrum to the same power-law function. The broad-line like structure down to $\\sim$ 4 keV is recognized both in the model and data, but this is naturally explained by  combination of the disk reflection, warm absorber and a mildly broad gaussian line that is not significantly red-shifted. We conclude  that an extremely broad emission line is not required if multiple warm absorbers are taken into account, which agrees with  Miller et al. (2008), while our model is simpler.   Also we point out that  normalization of the reflection component requires  $\\Omega/2\\pi \\sim 1$ in our model, not $>3$ as Miniutti et al. (2007)  advocates. Difference is mainly due to the fact that we included warm absorbers, and used neutral refection without relativistic smearing. So that the constancy of the reflection component is explained by the light-bending model, all the disk reflection {\\em must}\\/  take place within 3 $r_g$, in which case the reflection normalization requires $\\Omega/2\\pi \\sim 3$ (Miniutti \\& Fabian 2004). On the other hand, our result favors a smaller solid angle of the reflector. Also, invariability of the neutral reflection component suggests that  the reflection takes place far enough from the black hole, so that the intrinsic variation is smeared. If that is true, we would expect  a narrow line of the equivalent width $ \\sim 150$ eV (e.g., George and Fabian 1991) for $\\Omega/2\\pi \\sim 1$, much larger than the observed value ($\\sim$18 eV; Young et al.\\ 2005). If we allow the line width free, we obtain  the equivalent width $\\sim$ 110 eV (table \\ref{table0}), which may be reconciled with the reflector having $\\Omega/2\\pi \\sim$ 1. In this case, we require a mechanism to mildly broaden the line width up to $1 \\sigma \\approx 290$ eV. If this intrinsic line width is from Keplarian motion, with $v/c \\sim 0.29/6.35$, $v$ is $\\sim 14,000$ km/s, which does not seem too infeasible to be material on the innermost edge of the broad line region.   In summary, we found a characteristic spectral variation of MCG-6-30-15, that is explained by variation of ionization degree of the warm absorber accompanying the flux variation, while the neutral reflection component was invariable throughout the observation. Considering the warm absorber, the observed energy spectrum is explained by a mildly broadened iron line, which is not significantly red-shifted, and a neutral and constant disk reflection component of which solid angle is $\\Omega/2\\pi \\sim 1$. Still, there needs a mechanism to keep the reflection component constant while direct component is variable."
    },
    "0910/0910.5821_arXiv.txt": {
        "abstract": "In this letter, we report on dual-frequency European VLBI Network (EVN) observations of the faintest and least luminous radio cores in Seyfert nuclei, going to sub-mJy flux densities and radio luminosities around $10^{19}\\,$W\\,Hz$^{-1}$.  We detect radio emission from the nuclear region of four galaxies (NGC\\,4051, NGC\\,4388, NGC\\,4501, and NGC\\,5033), while one (NGC\\,5273) is undetected at the level of $\\sim 100\\, \\mu$Jy. The detected compact nuclei have rather different radio properties: spectral indices range from steep ($\\alpha>0.7$) to slightly inverted ($\\alpha = -0.1$), brightness temperatures vary from $T_B=10^5$ K to larger than $10^7$ K and cores are either extended or unresolved, in one case accompanied by lobe-like features (NGC\\,4051). In this sense, diverse underlying physical mechanisms can be at work in these objects: jet-base or outflow solutions are the most natural explanations in several cases; in the case of the undetected NGC\\,5273 nucleus, the presence of an advection-dominated accretion flow (ADAF) is consistent with the radio luminosity upper limit. ",
        "introduction": "Active Galactic Nuclei (AGN) are traditionally divided in radio quiet (RQ) and radio loud (RL). The latter are typically powerful radio sources ($P_r>10^{22}\\,$W\\,Hz$^{-1}$), with large scale radio lobes and bright compact cores; VLBI observations routinely target their nuclear regions, showing high brightness temperatures and in some cases jet knots with superluminal motions. Radio quiet AGN, such as Seyfert galaxies, are much fainter radio sources and their radio emission is confined to the sub-kpc scale. However, moderately deep VLA surveys show that most AGN are radio sources at some level \\citep[][~hereinafter HU01]{Nagar2002,Ho2001}.  While the origin of the radio emission in RL AGN is well established as synchrotron radiation from energetic particles in jets and lobes, the case for RQ AGN is much less clear.  Since nuclear structures in most Seyfert galaxies show complex morphologies, it is of fundamental importance to resolve them with high spatial resolution. Indeed, it has been shown that VLBI observations of the pc-scale and subpc-scale region of Seyfert nuclei and Low-luminosity AGNs (LLAGNs) are often successful in the determination of the physical parameters of the nuclear radio components (such as the brightness temperature, the spectral index, the jet motions, etc.) and therefore in the comprehension of the underlying physical mechanisms. For example, the VLBA study of the classical Seyfert 2 galaxy NGC\\,1068 has allowed the identification of the location of the hidden active nucleus and the attribution of the core radio emission to thermal free-free emission from an X-ray heated corona or wind arising from the disk \\citep{Gallimore2004}. In the case of the type 1 Seyfert NGC~4151 it has been possible to resolve the 0.2 pc two-sided base of a jet whose low speeds indicate non-relativistic jet motions, possibly due to thermal plasma \\citep{Ulvestad2005}. On the contrary, the VLBA analysis of the LLAGN NGC\\,4278 has shown that the radio emission of this source is emitted via synchrotron process by relativistic particles similarly to ordinary radio-loud AGN \\citep{Giroletti2005}. These studies indicate that the analysis of individual sources is very important for the determination of the different spatial components and physical parameters.  The five sources here presented belong to a complete distance limited ($d < 22\\,$Mpc) sample of 27 nearby Seyfert galaxies \\citep{Cappi2006}. Accurate multi-wavelength studies are available for this sample. In particular, VLA radio images at 1.4 and 5 GHz are presented in \\citetalias{Ho2001}, and VLA data at 15 GHz are also available for $25$ sample sources \\citep{Nagar2002}. VLBI observations have become available for a few sources through the years; $9$ sources of the sample have 5 GHz VLBA observations \\citep{Nagar2002}, while for the radio brightest galaxies, e.g.\\ NGC\\,1068 and NGC\\,3031, dedicated works are available \\citep{Gallimore2004,Bietenholz2004}. However, most of the weakest sources in the sample have never been observed with milliarcsecond resolution.  In the present letter, we report on new high sensitivity VLBI observations of five sources with $S\\sim1$ mJy, aimed at answering two fundamental questions: (1) how common/frequent is the presence of (sub-)parsec scale radio sources even in the weakest nuclei and (2) what do the properties of the detected radio sources tell us about the physics of individual sources and the viability of jet-base versus Advection-Dominated Accretion Flow \\citep[ADAF,][]{Narayan1994} explanations. The observations are detailed in \\S2 and the results presented in \\S3. A discussion and the main conclusions are given in \\S4 and \\S5, respectively. Throughout the paper, we define spectral indices such that $S(\\nu)\\propto \\nu^{-\\alpha}$ and adopt $H_0=70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_\\Lambda=0.73$ and $\\Omega_m=0.27$ \\citep{Spergel2003}. ",
        "conclusions": "With the VLBI observations presented in this letter, we have targeted the faintest and least luminous nuclei among well known local AGN, going to sub-mJy flux densities and radio luminosities around $10^{19}$ W Hz$^{-1}$.  These sources belong to a complete sample of nearby Seyfert galaxies, in which previous VLBI observations -- albeit limited to the brightest members -- had revealed an ubiquitous presence of sub-parsec cores and/or structures.  Similar findings had been reported on other relatively bright flat spectrum radio sources in LLAGN; for instance, \\citet{Nagar2002} found that almost all LLAGNs in the Palomar sample with S$_\\mathrm{VLA, 15\\, GHz}> 2.7$ mJy show mas-scale or sub mas-scale radio cores.  The 80\\% detection rate in the present sub-sample extends these findings to the lowest luminosity and flux density regimes.  While the presence of sub-pc scale radio emission appears thus to be ubiquitous, it is remarkable that it accounts only for a fraction between 5\\% and 40\\% of the emission detected on scales of a few tens of parsecs (eg.\\ as seen in VLA observations). As a consequence, a large fraction of the parsec scale radio luminosity is emitted in a diffuse region.  This could be the case of NGC\\,4388 and NGC\\,4501, in which we completely resolve the VLA sources at 5 GHz, but detect them at 1.6 GHz (size $\\ga 10$ mas). Similarly, the non detection of NGC\\,5273 requires that $>95\\%$ of the VLA 1.6 GHz flux density is emitted on scales larger than 20 mas (1.6 pc).  Even though much of the intermediate scale emission is resolved out, the majority of the nuclear regions do host compact radio sources at sub-parsec scales. In Table~\\ref{tabtot}, we report the multi-wavelength properties of the sources presented in this letter.  The common features in these nuclei are: (i) their extremely low radio luminosity (Cols.\\ 4-6), at the level of the least luminous Seyfert nuclei \\citep[such as the sub-mJy source NGC\\,4395,][]{Wrobel2006}; (ii) their extreme radio-quietness, for instance the ratios between the radio and nuclear X-ray emission (Col.\\ 8) are among the lowest ever measured for LLAGN \\citep{Panessa2007,Terashima2003}; and (iii) their low Eddington ratios (Col.\\ 9), a common trait in low luminosity nuclei \\citep{Ho2009}.  Radio emission from LLAGN cores is generally associated with accretion/ejection processes in the vicinity of a supermassive black hole \\citep{Falcke2000,Ulvestad2001,Nagar2002}. While some of the resolved extended emission could be of thermal origin, the small linear scales ($< 0.6$ pc) and comparatively high brightness temperatures ($T_B=10^{5-7}$ K) measured in our sources suggests that the weakest Seyfert nuclei could also be scaled-down versions of more luminous AGN. In this scenario, the origin of radio emission can be generally attributed to synchrotron emission from the base of a jet coupled with a low-power accretion disk \\citep{Falcke1999}.  Alternatively, the presence of an ADAF \\citep{Narayan1994} is invoked to explain the low luminosity of these sources, although the predicted radio emission often fails to reproduce the data, requiring combined jet/ADAF models \\citep[see e.g.,][]{Yuan2002}. On average, our results do not seem consistent with the presence of an ADAF alone, on the basis of the observed structures, sizes, flux densities, and spectral indices.  The two sources detected at both 1.6 and 5 GHz (NGC\\,4051 and NGC\\,5033) are most easily explained in terms of jet-base/outflow phenomena. First, the detection of three aligned sub-mJy components in NGC\\,4051 is indeed suggestive of ejection processes. The brightness temperature ($T = 2\\times10^5$ K) does not completely rule out a thermal origin, consistent with the presence of an outflow rather than a relativistic jet, similar to the case of NGC\\,4395, which shows an elongated low brightness temperature nuclear structure, tracing possible outflow emission \\citep{Wrobel2006,Christopoulou1997}.  The EVN radio core is also positionally coincident with a low luminosity H$_2$O maser, suggesting that the radio continuum may arise from the inner regions of a molecular disk or from a nuclear wind \\citep{Hagiwara2007}.  As for NGC\\,5033, its core presents all the hallmarks of a jet-base feature, being detected and unresolved at both at 1.6 and 5 GHz, with a brightness temperature lower limit of $T_B>1.3\\times10^7$ K, and a flat spectrum (consistent with $\\alpha=0$).  While these would somewhat be consistent with ADAF expectations, the 5 GHz luminosity is $\\sim$ 4 times in excess of what predicted on the basis of the observed L$_\\mathrm{2-10 keV}$ luminosity, implying that the ADAF model alone fails to account for the observed radio emission. The presence of a jet or an outflow component is therefore required and more consistent with the data. Indeed, the short jet-like extension found at larger scale with VLA \\citepalias{Ho2001,Perez-Torres2007} favours the jet-base hypothesis.  With all the caveats related to the unfavourable observing conditions at 5 GHz, the two sources detected only at 1.6 GHz (NGC\\,4388 and NGC\\,4501) seem to be also at odds with an ADAF scenario. The non detection at high frequency points to a steep spectral index and/or to a rather extended structure, at scales of $10^6 R_S$ or more, in contrast with the flat/inverted spectrum and the compact structure predicted by ADAF models. Although consistent with an angular scale of several mas, thermal emission is also unlikely given the $\\ga 10^6$ K brightness temperatures observed at 1.6 GHz. NGC\\,4388 and NGC\\,4501 are also the only type 1.9 Seyferts in our sample, showing heavily absorbed/weak X-ray spectra \\citep{Cappi2006}.  Interestingly, the only undetected source in our sample is NGC\\,5273, a type 1.5 Seyfert, in which the nucleus is seen directly, displaying broad emission lines and a bright X-ray spectrum \\citep{Cappi2006}. Either the VLA emission is resolved out or the source is variable. Indeed, the source was initially not revealed at 8 GHz with the VLA \\citep[$S < 0.23$ mJy,][]{Kukula1995} and lately detected with $S = 0.6$ mJy by \\citet{Nagar1999}. The comparison between the X-ray and the EVN radio luminosity upper limit reveals that this source is extremely radio quiet. Indeed, this is the only case in which the radio data are consistent with a pure ADAF accretion mechanism, since our upper limit for this nucleus is 4 times above the radio core luminosity derived from the observed X-rays \\citep{Yi1998}."
    },
    "0910/0910.1069_arXiv.txt": {
        "abstract": "From multi-epoch adaptive optics imaging and integral field unit spectroscopy we report the discovery of an expanding and narrowly confined bipolar shell surrounding the helium nova V445 Puppis (Nova Puppis 2000). An equatorial dust disc obscures the nova remnant, and the outflow is characterised by a large polar outflow velocity of 6720 $\\pm$ 650 km s$^{-1}$ and knots moving at even larger velocities of 8450 $\\pm$ 570 km s$^{-1}$. We derive an expansion parallax distance of 8.2 $\\pm$ 0.5 kpc and deduce a pre-outburst luminosity of the underlying binary of $\\log L/L_{\\odot} = 4.34 \\pm 0.36$. The derived luminosity suggests that V445 Puppis probably contains a massive white dwarf accreting at high rate from a helium star companion making it part of a population of binary stars that potentially lead to supernova Ia explosions due to accumulation of helium-rich material on the surface of a massive white dwarf. ",
        "introduction": "Accretion onto compact objects can lead to a range of explosive phenomena since their accreted layers can reach densities and temperatures high enough to initiate nuclear burning. Helium novae are expected to occur during periods of high mass-accretion rates ($\\dot{M} \\sim 10^{-9} - 10^{-6}$ M$_{\\odot}$ yr$^{-1}$) of He-rich material through unstable helium burning via helium shell flashes \\citep{iben91}. In such events, typically $\\sim 10^{-4}- 10^{-2}$ M$_{\\odot}$ of fuel is ignited on the surface of the accretor \\citep{kato99,yoon04,bild07}. Models of the binaries that  experience such helium novae generally consist of a 0.6 -- 0.8 M$_{\\odot}$ CO white dwarf (the accretor) with a lower-mass H-deficient companion (the donor). The donor can be either a non-degenerate helium star \\citep[e.g.,][]{yoon04} or a semi-/fully degenerate star. The latter are usually referred to as AM CVn systems, representing a class of ultra-compact helium-transferring binaries \\citep{nele04,bild07,roel07}.  Some of these compact binaries may lead to sufficient mass accumulation onto the primary to be viable supernova Ia progenitors after successive He novae \\citep{iben94,kato89,yoon03,tutu07,wang09}. Accreting massive white dwarfs in close binary systems are the favoured progenitors of supernova Ia explosions \\citep{branch07,par07}. Possible binary  channels are typically divided into single-degenerate scenarios involving hydrogen-rich companions versus double-degenerate progenitors. The latter include the white dwarf plus He-star and double white dwarf pathways. To date, only about a dozen promising single-degenerate type Ia progenitors have been identified \\citep[see, e.g.,][]{par07}. U~Scorpii and RS Ophiuchi represent the two different classes of hydrogen-rich companion stars in the supernova Ia progenitors: main-sequence or slightly evolved stars, and red giants, respectively \\citep{li97}. Ongoing supernova searches are, however, revealing a growing diversity of type Ia explosion events, which challenge current formation scenarios. It could be that, despite their low predicted frequency, the helium star donor channel to type Ia explosion is successful in explaining some of the diversity observed, such as the population of young Ia progenitors \\citep{manu06,wang09}.  \\object{V445 Puppis} (Nova Puppis 2000) is the first, and so far only, helium nova detected \\citep{ashok03,kato03}. The nova outburst of V445 Puppis was first noticed on 23 November 2000 \\citep{kato01}, and the optical outburst characteristics are those of a slow nova. Optical \\citep{wagn01,iijim08}, near-infrared \\citep{ashok03} and mid-infrared \\citep{lynch01} spectra of V445 Puppis obtained during outburst immediately revealed its most peculiar and defining characteristic: the complete absence of hydrogen in the ejecta.  The outburst light curve has been modelled \\citep{kato03,kato08} by free-free emission and an optically thick wind, following a helium shell flash on the surface of an accreting massive white dwarf. A lower limit on the mass of the accreting white dwarf of $M$ $\\ge$ 1.35 M$_{\\odot}$ is inferred by \\citet{kato08}, making V445 Puppis a possible progenitor of a supernova Ia from a helium-rich donor channel. In these models, the white dwarf is expected to grow in mass as the mass ejected through repeated helium nova outbursts is less than the accreted mass \\citep{kato99}.  V445 Puppis provides the first empirical benchmark for a helium-dominated outburst on the surface of a white dwarf. In Section 2 we present new observations of V445 Puppis taken 5 -- 7 years after outburst, and determine the distinct evolution of the resolved nova shell from a spatio-kinematic analysis (Section 3). This leads to a robust distance determination to V445 Puppis. In Section 4, we discuss the implications of the distance derived here on the nature of the underlying binary. ",
        "conclusions": "\\subsection{Pre-outburst conditions}  With the distance known, the nature of the helium nova progenitor can be constrained from pre-outburst observations. Unfortunately, not much is known of V445 Puppis prior to outburst. Apart from optical ($V = 14.5$ mag, from VSNET, see \\citet{ashok03}) and near-infrared (2MASS: $K_s = 11.52$ mag) flux measurements (see lower panel of Fig.~\\ref{woudtfig1}), neither spectrum nor orbital period have been obtained prior to outburst. We verified the plate archives at the Harvard-Smithsonian Center for Astrophysics for prior outbursts (or sustained long periods of obscuration indicative of a missed outburst). Despite good time coverage and the fact that V445 Puppis could be identified on many plates at approximately constant brightness (based on visual comparision with nearby stars in the field), we could find no signatures of a previous outburst in the 1897 -- 1955 time frame.  V445 Puppis is located at the low Galactic latitude of $b = -2.19^{\\circ}$, which at a distance of 8.2 $\\pm$ 0.5 kpc translates to a height below the Galactic Plane of 313 $\\pm$ 19 pc. At that distance, and that far below the Galactic Plane, one can use the IRAS/DIRBE Galactic reddening maps \\citep{schleg98} to gauge the Galactic reddening towards V445 Puppis. In Fig.~\\ref{woudtfig6} we show the ratio of the measured reddening (at a given distance) to the total line-of-sight Galactic reddening for open clusters in the Milky Way \\citep{khar05} in the Galactic latitude range $1^{\\circ} \\le |b| \\le 4^{\\circ}$ and for $E(B-V)_{\\rm Schlegel} \\le 2.5$, as a function of height above/below the Plane and distance, respectively. The reddening maps of \\citet{schleg98} are not well-calibrated at low Galactic latitude; from colors of galaxies discovered at low Galactic latitude behind the southern Milky Way it appears that the reddening values from Schlegel et al.~are too high \\citep{schroed07} and that the true value is 87\\% of the \\citet{schleg98} reddening, e.g. $E(B-V)^{\\rm cal}_{\\rm Schlegel} = 0.87 E(B-V)_{\\rm Schlegel}$.  Comparing the location of V445 Puppis to open clusters within 10 degrees of V445 Puppis (big filled circles in Fig.~\\ref{woudtfig6}), we deduce that Galactic foreground reddening towards V445 Puppis is probably in the range of $E(B-V) = 0.51 - 0.68 = 0.75 - 1.0 \\, E(B-V)^{\\rm cal}_{\\rm Schlegel}$, as marked by the arrow in Fig.~\\ref{woudtfig6}. The lower limit given here could be a conservative one due to a selection effect; given the relative poor angular resolution of the reddening maps, the open clusters which set the lower ratio limits (those observed at large distances and large height above/below the Galactic plane) are likely to be observed through small-scale patches of relatively lower extinction compared to their surrounding.  Finally, taking the equivalent widths of the two interstellar Na D components \\citep{iijim08}, a value of $E(B-V) = 0.62$ is inferred using the calibration of \\citet{mun97}. This value falls within our deduced limits.  Taking the lower limit of the Galactic reddening towards V445 Puppis ($E(B-V) = 0.51$), we derive a pre-outburst color of V445 Puppis of $(V-K)^0 = 1.58$ mag, assuming a standard Galactic extinction law \\citep{card89}. This is significantly redder than the likely binary progenitors -- high $\\dot{M}$ AM CVn systems or systems with helium stars donors -- which typically have $(V-K)^0 \\approx -0.6$. A red giant donor (symbiotic nova) can be excluded on the basis of the pre-outburst colors.  This suggests the presence of substantial circumstellar (CS) reddening before the November 2000 outburst, which is not unreasonable given the possibility of previous outbursts or material expelled during a common envelope phase. The CS reddening around V445 Puppis does not necessarily have to follow a standard Galactic extinction given the current carbon-rich ejecta; \\citet{berge99} find some deviations from a standard reddening law for CS reddening around a number of carbon-rich R CrB stars. In the right panel of Fig.~\\ref{woudtfig3}, we compare the line profile of the 7065 {\\AA} and 2.0581 $\\mu$m He\\,I recombination lines, obtained close in time and both normalised by peak blueshifted emission. As expected, the redshifted component of the optical line (at +1140 km s$^{-1}$) appears dimmer compared to its near-infrared counterpart, due to dust obscuration within the shell. For a standard Galactic reddening law \\citep{card89}, we expect a ratio of peak emission of near-infrared (2.0581 $\\mu$m) to optical (7065 {\\AA}) of 7. The ratio observed in V445 Puppis is 3.4, substantially less and indicative of a low ratio of total-to-selective extinction, $R_V \\approx 2.5$ \\citep{fitzp99}. The shell of V445 Puppis offers an opportunity to determine the nature of the dust extinction in carbon-rich outflows. The new X-shooter instrument on the VLT will be ideal to obtain simultaneous optical to near-infrared medium-to-high resolution spectroscopy of both NE and SW shells.  To make the pre-outburst color of V445 Puppis consistent with the intrinsic colors of likely progenitor binaries, a (maximum) additional color excess of $A_V - A_K$ = 2.2 mag requires $A_V$ = 2.5 mag for $R_V = 3.1$ (standard reddening law), or $A_V$ = 2.8 mag for $R_V = 2.5$.  The uncertainty in $A_V$ introduced by the different reddening laws, however, is small compared to the uncertainty in the bolometric correction needed to arrive at the pre-outburst luminosity of V445 Puppis.  We thus determine a pre-outburst extinction-corrected brightness of at least $V^0 = 12.9$ mag (corrected for Galactic reddening only), but more likely $V^{0} \\approx 10.1 - 10.4$ mag (including circumstellar reddening correction), with the range in the latter allowing for differences in the CS reddening law. The corresponding absolute magnitudes are $M_V^{0} = -1.7$ and $-4.2$ to $-4.5$ mag, respectively. For a range of likely bolometric corrections (BC = $-1$ to $-2.5$, \\citet{roel07,heb81}), the pre-outburst luminosity of V445 Puppis is $\\log L/L_{\\odot} = 3.3 \\pm 0.3$ (corrected for Galactic extinction only), or $\\log L/L_{\\odot} = 4.34 \\pm 0.36$ (including circumstellar reddening correction).  It has not escaped our attention that hydrogen-deficient carbon (HDC) stars also occupy this luminosity regime. With a larger distance -- this paper -- and a larger Galactic reddening correction \\citep{iijim08}, the possibility that V445 Puppis belongs to the class of HDC stars can no longer be rejected on luminosity grounds, cf. \\citet{ashok03}. The (Galactic) extinction-corrected $(V-K_S)^0$ colors of V445 Puppis are not too dissimilar from the HDC star HD\\,182040 \\citep{brun98}. It is not clear, though, what can cause a luminosity increase of 6 magnitudes and a rapidly expanding shell from a HDC star; the overall light curve appears much better described by an event near a compact, massive white dwarf.  \\subsection{Shaping of the nebula}  Deviations from the simple outflow model presented in Section 3 could provide clues to the mechanism responsible for shaping the nebula. Although the dynamical model fits the shell remarkably well, the largest deviations occur nearest to the nova remnant at the earliest epochs. The nebula initially had a very narrow waist, followed by a rapid broadening of the waist close to the nova remnant.  To produce very narrow waists in PNe, collimated fast winds (CFW) are needed in systems with high density gas in an equatorial plane close to the source of the CFW \\citep{sok00}. It is unclear what collimated the fast wind in V445 Puppis. At the moment it seems likely that an equatorial disk/torus already existed pre-outburst, given the large pre-outburst circumstellar reddening deduced in Section 4.1. Once the present optically-thick dust disk clears (Fig.~\\ref{woudtfig2} -- presumably additional material was ejected in the equatorial plane during the November 2000 outburst) and a relatively unobscured view of the nova remnant is possible, alternative origins of the collimated outflow could be investigated, e.g. the presence of a rapidly rotating magnetic white dwarf.  The shell of V445 Puppis is unique for a classical nova. Although many novae show bipolar shells, for example the classical nova HR Del \\citep{har03} and recurrent nova RS Oph \\citep{bode07,sok08}, the shell in V445 Puppis reflects the most collimated outflow seen in any nova. The knots are moving at jet-like velocities of $\\ga$ 8000 km s$^{-1}$. They can be a consequence of a jet-activity about 350 days after the main outburst which is not long after the sudden drop off in the $V$ band is seen \\citep{ashok03} and coincides with the radio flare \\citep{rup01}. In the symbiotic nova CH Cygni, a drop in $V$ band magnitude is associated with the onset of a radio jet ejection, see e.g. \\citet{karov07}.  \\subsection{V445 Puppis as a proto-type helium nova}  At a distance of 8.2 kpc, the maximum brightness at outburst ($V = 8.6$ mag) is consistent with the Eddington luminosity of a massive white dwarf. Moreover, the pre-outburst luminosity of V445 Puppis of $\\log L/L_{\\odot} = 4.34 \\pm 0.36$ appears to rule out an AM CVn progenitor and may require both a luminous accretion flow as well as a bright donor. For comparison, a 1.3-M$_{\\odot}$ white dwarf accreting at $\\dot{M} \\sim 10^{-6}$ M$_{\\odot}$ yr$^{-1}$ would give an expected accretion luminosity of $\\log L/L_{\\odot} = 3.7$. Similarly, the luminosity of a moderately massive, evolved helium star is around $\\log L/L_{\\odot} \\approx 3.7$ \\citep{yoon03}.  With the distance determined in this paper, a revised mass of the dust shell of $1.5 \\times 10^{-5}$ M$_{\\odot}$ is derived following \\citet{lynch04}. This is consistent with \\citet{kato03}, but we note that the mass depends strongly on the assumed temperature of the shell and will ultimately be best constrained by mid- to far-infrared data of V445 Puppis recently obtained by {\\it Spitzer}. This mass is likely to be a lower limit given the large optical depth of the equatorial dust disk.  The various components of V445 Puppis (the dust disk, the bipolar shell) clearly show that the immediate environment of V445 Puppis is sculpted either directly or indirectly by (repeated) outbursts. If V445 Puppis is representative of its class (of helium novae), the helium counterparts to the classical hydrogen-rich novae result in more substantial circumbinary reddening due to the carbon-rich outflow.  Validation of the white dwarf + helium star model as the appropriate binary configuration of V445 Puppis will come from the determination of its orbital period. Unfortunately, this can only be done once the optically thick dust disk surrounding the nova remnant becomes transparent. Given the low inclination of the bipolar outflow, the implied inclination of the equatorial plane which corresponds to the orbital plane of the binary is $\\sim 86^{\\circ}$.  We thus expect this to be an eclipsing binary, which would further facilitate our ability to constrain the component masses.  \\subsection{V445 Puppis as a candidate Ia progenitor}  There are several indirect suggestions that the white dwarf in V445 Puppis is massive: the maximum luminosity during outburst, the large velocities of the ejected blobs, and its pre-outburst (accretion) luminosity.  The lack of strong soft X-ray emission pre-outburst, as is expected from a supersoft source \\citep{yoon03} could be explained by the substantial interstellar and circumstellar absorption which currently totally obscures the underlying accreting binary. V445 Puppis has not been detected in the ROSAT all-sky survey.  Our derived luminosity is consistent with the massive white dwarf + helium star model of \\citet{yoon03}, but see also \\citet{kato08}. It is likely that V445 Puppis belongs to a class of binaries that, in principle, could lead towards supernovae of type Ia. Given that the expected lifetime of a massive white dwarf + helium star binary is short, these progenitors are associated with the relatively young ($\\sim 10^8$ yr) population of Ia SNe.  Whether V445 Puppis itself eventually leads to a supernova Ia event, or if the current helium nova has pre-empted that pathway, depends critically on the mass accumulation efficiency on the surface of the white dwarf and the mass ejected during this nova outburst. This remains a topic of debate."
    },
    "0910/0910.4961_arXiv.txt": {
        "abstract": "As the Galaxy evolves, the abundance of deuterium in the interstellar medium (ISM) decreases from its primordial value: deuterium is ``astrated\".  The deuterium astration factor, $f_{\\rm D}$, the ratio of the primordial D abundance (the D to H ratio by number) to the ISM D abundance, is determined by the competition between stellar destruction and infall, providing a constraint on models of the chemical evolution of the Galaxy.  Although conventional wisdom suggests that the local ISM (i.e., within $\\sim 1-2$~kpc of the Sun) should be well mixed and homogenized on timescales short compared to the chemical evolution timescale, the data reveal gas phase variations in the deuterium, iron, and other metal abundances as large as factors of $\\sim 4-5$ or more, complicating the estimate of the ``true\" ISM D abundance and of the deuterium astration factor. Here, assuming that the variations in the observationally inferred ISM D abundances result entirely from the depletion of D onto dust, rather than from unmixed accretion of nearly primordial material, a model-independent, Bayesian approach is used to determine the undepleted abundance of deuterium in the ISM (or, a lower limit to it).  We find the best estimate for the undepleted, ISM deuterium abundance to be (D/H)$_{\\rm ISM}\\geq (2.0 \\pm 0.1) \\times 10^{-5}$.  This result is used to provide an estimate of (or, an upper bound to) the deuterium astration factor, $f_{\\rm D} \\equiv$~(D/H)$_{\\rm P}/$(D/H)$_{\\rm ISM} \\leq 1.4 \\pm 0.1$. ",
        "introduction": "Deuterium is created in an astrophysically interesting abundance only during big bang nucleosynthesis (BBN) \\citep{boes,steigman07}, after which, in the post-BBN Universe, its abundance ($y_{\\rm D} \\equiv 10^5 (\\rm D/H)$) decreases monotonically due to the processing of gas through succeeding generations of stars where deuterium is completely destroyed \\citep{els,pf03}. Consequently, deuterium plays a special role in cosmology, nuclear astrophysics, and in Galactic chemical evolution \\citep{ytsso,boes,st92,st95,vangioni,tosi96,schramm,romano06,srt,steigman07,pf08}. Its relatively simple Galactic evolution permits us to use deuterium to determine the fraction of interstellar gas that has been processed through stars \\citep{st92,st95}.  By comparing the primordial and Galactic deuterium abundances one can learn about and discriminate among different Galactic chemical evolution models \\citep{vangioni,tosi96,romano06,srt,pf08}.  Together with non-BBN constraints on the baryon (nucleon) density from observations of the cosmic microwave background \\citep{spergel,Dunkley09,komatsu09,komatsu10}, the primordial deuterium abundance, \\ydp, is predicted by BBN \\citep{cyburt03,coc,steigman07,vs,cyburt08}.  The predicted abundance is in excellent agreement with observations of deuterium in high-redshift, low-metallicity QSO absorption line systems (QSOALS) \\citep{omeara,pettini}.  Since deuterium is destroyed in the Galaxy as gas is cycled through stars, (D/H)$_{\\rm ISM} \\leq$ (D/H)$_{\\rm P}$.  However, all successful chemical evolution models require infall to the disk of the Galaxy of unprocessed (or, nearly unprocessed) gas \\citep{tosi96,pf08}, and such deuterium-rich (and metal-poor) gas would raise the ISM D/H ratio closer to the primordial value.  As a result, comparison of the primordial and ISM deuterium abundances provides a constraint on infall and, therefore, on chemical evolution models.  Observations over the past decade and more of the deuterium abundance in the relatively local interstellar medium (ISM) reveal an unexpectedly large scatter in D/H \\citep{jenkins,sonneborn,hebrard,hoopes}, challenging the conventional wisdom of a well mixed ISM.  For example, as shown in Figure~\\ref{fig:lgyvslgh}, the absorption line measurements from the {\\it Far Ultraviolet Spectroscopic Explorer} (FUSE) reveal variations of a factor of $\\sim 4$ ($0.5 \\la y_{\\rm D,ISM} \\equiv 10^{5}$(D/H)$_{\\rm ISM} \\la 2.2$) in the {\\em gas-phase} D/H ratios over lines of sight (LOS) to background stars within $\\sim 1-2$ kpc of the Sun.  Moreover, the variations in the observed, gas phase D/H abundances are found to correlate {\\em positively} with the abundances of refractory elements such as Ti \\citep{prochaska,lallement08}, and a similar correlation is also found between D and other metals such as Fe and O \\citep{linsky,srt}.  However, this positive correlation between the abundances of D and the metals is opposite to the trend expected from stellar nucleosynthesis (D decreasing as the metals increase).  Motivated by the observed correlations and by the very large spread in the D/H ratios inferred from the FUSE and other data sets, it has been proposed that the large variations in {\\em gas-phase} D/H may be due to the depletion of gas phase deuterium onto dust grains \\citep{jura,draine04,draine06}.  If this is the case, then the FUSE (and other) gas-phase absorption-line measurements reveal that depletion has not been well mixed (homogenized) in the ISM and, the data may only provide a {\\em lower} limit to the {\\em true} (undepleted) ISM deuterium abundance and, therefore, only an {\\em upper} limit to the deuterium astration factor, $f_{\\rm D} \\equiv y_{\\rm DP}/y_{\\rm D,ISM}$.  \\begin{figure}[t] \\epsscale{0.4} \\plotone{lgyvslgh1.eps} \\caption{The logs of the deuterium abundances versus the logs of the \\hi~column densities [cm$^{-2}$] for the 49 FUSE LOS (see the text).  The filled symbols are for the 41 LOS which have iron abundance data, while the open symbols are for the 8 LOS which lack iron abundances.  The squares (blue) are for the LOS within the Local Bubble (LB) and the circles (red) are for the non-LB (nLB) LOS.  The solid line is at the mean D abundance for the LB LOS (log $y_{\\rm D,LB} = 0.19$); the dashed line is its extension to the nLB LOS.} \\label{fig:lgyvslgh} \\end{figure}  In Figure~\\ref{fig:lgyvslgh} the logs of the deuterium abundances (log~$y_{\\rm D} \\equiv 5~+~$logN(\\di) $-$ logN(\\hi)) are shown as a function of the logs of the \\hi~column densities for the 49 LOS with \\di~column densities from Table 2 of \\citet{linsky}, supplemented by data from \\cite{oliveira2006} and \\cite{dupuis2009}.  Of these 49 LOS, 41 have iron abundance measurements (filled symbols); 21 of the 49 LOS are within the Local Bubble (LB; see \\citet{linsky} for a discussion of the LB); the remaining 28 non-Local Bubble (nLB) LOS are toward stars beyond the LB (see \\S2.1).  The unexpectedly large spread among the observationally inferred ISM D abundances complicates any estimate of \\ydism.  Recognizing this point, \\citet{linsky} chose for their estimate of a {\\em lower bound} to the true (undepleted) ISM deuterium abundance the mean of the five {\\em highest} D/H ratios finding \\ydism~$\\geq 2.17 \\pm 0.17$ or, when including corrections (see the discussion in \\S2.1) for N(\\hi) and N(\\di) for those lines of sight outside of the Local Bubble, \\ydism~$\\geq 2.31 \\pm 0.24$.  These estimates are quite close to the {\\it lower bound} to the primordial abundance estimated from the QSOALS, $y_{\\rm DP} = 2.82^{+0.20}_{-0.19}$ \\citep{pettini}\\footnote{ Since this estimate of \\ydp~relies on only 7 high redshift, low metallicity LOS and, since the dispersion in deuterium abundances among them is unexpectedly large, some prefer to adopt a so-called ``WMAP D abundance\".  WMAP does not observe deuterium.  Rather, the WMAP-determined baryon density parameter may be used in a BBN code to {\\it predict} the relic D abundance.  If the \\cite{komatsu10} estimate of the baryon density is adopted, the BBN predicted primordial D abundance is \\ydp~$= 2.5 \\pm 0.1$.  However, the WMAP collaboration also provides an estimate of the effective number of neutrinos, \\citep{komatsu10} which, when used along with their baryon density estimate in a BBN code, leads to a {\\it different} predicted primordial D abundance \\ydp~$= 3.0 \\pm 0.4$.  The difference between these two predictions reflects the difference between the standard model effective number of neutrinos expected and the WMAP observed value, which is within $\\sim 1.5\\sigma$ of expectations.  So, which, if either, ``WMAP D abundance\" is preferred?  Here, we compare observations to observations, not to model-dependent predictions and, we adopt the \\cite{pettini} value for \\ydp.}, suggesting a small upper bound to the D astration factor, $f_{\\rm D} \\leq 1.30^{+0.14}_ {-0.13}$ or, an even smaller value, $f_{\\rm D} \\leq 1.22 \\pm 0.15$ for the LB-corrected, nLB deuterium abundances.  To account for such a high ISM deuterium abundance and such a small D astration factor, a very high infall rate of pristine material would be needed, challenging many Galactic chemical evolution models \\citep{romano06,pf08}.  By limiting themselves to the five highest D abundances, \\citet{linsky} ignore the lower deuterium abundances along many more LOS which are consistent with them within the errors, potentially biasing their estimates of \\ydism~and of $f_{\\rm D}$.  In an attempt to address this issue, \\citet{srt} (SRT), used the 18 highest D/H ratios from the FUSE data (see Table 3 of \\citet{linsky}), finding an ISM D abundance of $y_{\\rm D,ISM} = 1.88 \\pm 0.11$, corresponding to a D-astration factor $f_{\\rm D} \\leq 1.50^{+0.14}_{-0.13}$, consistent with at least some of the otherwise successful chemical evolution models discussed in SRT.  In fact, the data in Table 2 of \\citet{linsky} reveal that there are 19 LOS with central values of log \\yd~$\\geq 0.20$. The weighted mean (along with the error in the mean) for these 19 D abundances is log $y_{\\rm D,19} = 0.26 \\pm 0.01$, corresponding to $y_{\\rm D,19} = 1.8 \\pm 0.1$.  For these 19 LOS the reduced $\\chi^{2} = 0.85$, confirms that the weighted mean provides a good description of their D abundances.  As more LOS with lower D abundances are added, the weighted mean decreases, but the reduced $\\chi^{2}$ increases, so that \\ydism~$\\ga 1.8 \\pm 0.1$ is likely a robust {\\it lower} bound to the ISM D abundance and, $f_{\\rm D} \\la 1.5 \\pm 0.1$ (log $f_{\\rm D} \\la 0.19 \\pm 0.03$) is a robust {\\it upper} bound to the deuterium astration factor.  Surely, there must be a better way to find a reliable estimate of the maximum value of the deuterium abundance in the local ISM while accounting for the observational errors in the individual D abundance determinations.  In \\S 2 we describe a Bayesian analysis designed to find the best estimate of the maximum, gas phase (undepleted), ISM deuterium abundance, \\ydmax, from data with non-negligible errors and we apply it to the FUSE data set.  On the assumption that the spread in the observed D abundances is the result of incompletely homogenized dust depletion, \\ydism~$\\geq$ \\ydmax~and $f_{\\rm D} \\equiv y_{\\rm DP}/y_{\\rm D,ISM} \\leq y_{\\rm DP}/y_{\\rm D,max}$.  Our results are summarized and our conclusions presented in \\S 3.  As already noted, the dust depletion hypothesis suggests that there should be a correlation between the D and Fe abundances. Although this trend appears to be present at some level \\citep{linsky}, a deeper analysis suggests that the simplest interpretation of the depletion hypothesis may need to be modified~\\citep{srt}. For example, since D is likely attached to the mantles of dust grains \\citep{tielens}, while Fe is mostly contained in the cores of the grains, D will be removed from grains more easily than Fe. Thus, although strong shocks may destroy grains completely, weak shocks may only remove D from grains while Fe might stay locked within their cores \\citep{draine79}.  This effect then may explain the scatter observed in the relation between the D and Fe depletions.  Thus, in a follow-up paper in preparation, the consistency of the depletion hypothesis with the FUSE data for D and Fe (and O) is investigated, and some possible modifications to its simplest version are explored \\citep{sp}. ",
        "conclusions": "In the decades prior to the current era of precision cosmology the primordial abundance of deuterium provided the only quantitative cosmological baryometer~\\citep{boes}.  Although, at present, the deuterium abundance is only measured along seven, high-redshift, low-metallicity LOS to background quasars~\\citep{omeara,pettini}, the inferred primordial D abundance, \\ydp~= $2.8  \\pm 0.2$, is in excellent agreement with the non-BBN inferred baryon density parameter~\\citep{steigman07}.  In the post-BBN Universe deuterium is destroyed as gas is cycled through stars, so that comparing the abundance of deuterium in the ISM of the Galaxy with the primordial D abundance provides an estimate of the virgin fraction of the ISM (\\ie the amount of gas presently in the ISM which has never been cycled through stars), constraining models of Galactic chemical evolution~\\citep{st92,st95,srt,pf08}.  According to conventional wisdom the deuterium-free, metal-enhanced products of stellar nucleosynthesis should be well-mixed in the local ISM. In contrast, the FUSE data on the abundances of deuterium and several metals (\\eg iron, oxygen, etc.) along LOS within $\\sim 1 - 2$ kpc of the Sun reveal a much different picture.  The FUSE~\\citep{linsky} and earlier observations \\citep{jenkins,sonneborn,hebrard,hoopes,prochaska} reveal unexpectedly large gas phase variations in \\yd~(and in the abundances of iron, oxygen, etc.) within the local ISM, as shown for FUSE deuterium data in Figure~\\ref{fig:lgyvslgh}.  It has been proposed that the large variations observed in the local ISM D abundances can be accounted for by preferential depletion of deuterium (relative to hydrogen) onto dust \\citep{jura,draine04,draine06}, although incompletely mixed infall of relatively unprocessed, deuterium-enhanced, metal-free material may have contributed to some of the observed variations~\\citep{st92,st95,srt}.  The large variations among the ISM D abundances, along with observational errors and the possible contributions from dust depletion and infall, complicate using the D observations to provide a robust estimate of the ISM D abundance which, in combination with the primordial D value, can lead to a constraint on the deuterium astration factor, $f_{\\rm D}$.  The key question is, given the data (with its errors), how to find the best estimate of the ``true\", undepleted, ISM D abundance?  Here, to address this question, the limits to the true, undepleted, ISM D abundance were investigated employing a model-independent Bayesian statistical analysis similar to that used by \\citet{hogan} to infer the primordial helium abundance from a set of helium abundance observations.   It was assumed, along with \\citet{linsky}, that the spread in the observed D abundances is the result of incompletely homogenized D depletion onto dust in the local ISM. In our analysis this is modeled by five different probability distributions (priors) for the \\yd~values.  The \\yd~(actually, log~\\yd)~values shown in Figure \\ref{fig:lgyvslgh} suggest that, given the relatively large errors, a uniform (top-hat) distribution, favoring neither low-D nor high-D may be a good approximation to the data.  To explore the sensitivity of our result to the choice of the prior, we first considered two asymmetric distributions -- a positive-bias  prior favoring low depletion, and a negative-bias prior favoring large depletion, as well as two other priors -- an M-shaped distribution favoring both low and high depletion, and a complementary, $\\Lambda$-shaped distribution.  Using the FUSE deuterium observations along all 49 LOS \\citep{linsky,oliveira2006,dupuis2009}, we found the likelihoods in the \\{\\ydmax,$w$\\} plane for the five choices of the Bayesian priors (see Figures~\\ref{fig:allallpriors} and \\ref{fig:lbnlballM}).  For all priors, the Bayesian analysis of the full data set requires significant depletion (\\eg $w \\neq 0$ at greater than 99.9\\% confidence).  Comparing the maximum likelihood values for the five different distributions, we find that the bimodal, M-shaped distribution provides the best fit to the observed data (see Table 1).  However, it is important to notice that the shapes of the priors require an additional assumption in our analysis, so that the M-shaped distribution is the most model-dependent. \\footnote{The M-shaped prior favors both low and high levels of D depletion while strongly disfavoring intermediate depletion, suggesting that two competing processes may be at work: depletion onto dust and evaporation from dust, perhaps due to exposure to shocks.  To fit the M-prior scenario, both processes would have to be efficient and rapid to account for the deficit of intermediate D abundances.  The distribution of the presently available data (with its errors) is inconclusive and does not strongly favor any of the adopted prior distributions.  When more data become available, the Bayesian approach presented here may be used to learn more about the mechanism of deuterium depletion onto dust.}  In contrast, the top-hat prior is the least model-dependent, favoring all levels of depletion equally.  Given our ignorance of the detailed depletion mechanisms responsible for the observed scatter in the gas phase ISM deuterium abundances, we prefer to adopt for our estimate of the undepleted, ISM deuterium abundance, the result of the simplest, top-hat prior, whose maximum likelihood is similar to that of the best-fitting M-distribution, \\beq y_{\\rm D,ISM} \\geq y_{\\rm D,max} = 2.0 \\pm 0.1 = 2.0(1 \\pm 0.05) \\eeq This value is our best estimate of the true ISM D abundance based on the available deuterium observations in the local ISM and is independent of any model-dependent assumptions about galactic chemical evolution.  Combining our result with \\ydp~= $2.8 \\pm 0.2 = 2.8(1 \\pm 0.07)$ \\citep{pettini} (which, recall, provides a {\\it lower} bound to the primordial abundance), yields a limit to the deuterium astration factor $f_{\\rm D} \\leq 1.4 \\pm 0.1$ (for the M-prior, $f_{\\rm D} \\leq 1.6 \\pm 0.1$), consistent with most, but not all, Galactic chemical evolution models \\citep{srt,pf08,romano09}.  If, on the other hand we compared this \\ydism~value to the BBN + WMAP inferred primordial D abundance, \\eg \\ydp~= $2.5 \\pm 0.1$ \\citep{steigman10}, \\ydp~= $2.5 \\pm 0.2$ \\citep{cyburt08} or, the prediction inferred when including the WMAP-determined effective number of neutrino species \\citep{komatsu10}, \\ydp~= $3.0 \\pm 0.4$, the resulting deuterium astration factor would be somewhat lower in first two cases, $f_{\\rm D} \\approx 1.3 \\pm 0.1$, which is marginally problematic for some GCE models.  In contrast, for the \\citet{komatsu10} value of \\ydp, $f_{\\rm D} \\leq 1.5 \\pm 0.2$, which is entirely consistent with GCE models.  As seen in Figures~\\ref{fig:lgyvslgh} -- \\ref{fig:lbnlball}, for the LB there is little scatter among the gas-phase D abundances.  The small scatter is entirely consistent with the observational errors ($w = 0$) and all LB D abundances are consistent, within the errors, with \\ydlb~= $1.5(1 \\pm 0.03)$.  This suggests that for the Local Bubble, \\ydism~$\\geq$ \\ydlb~and, $f_{\\rm D,LB} \\leq 1.8 \\pm 0.1$, consistent with all the successful chemical evolution models identified in SRT.  However, while the uniform LB D abundance suggests that D may be {\\it undepleted} in the LB, for all LOS, \\ydmax~$\\approx 1.3$ \\ydlb, suggesting either that D is {\\it depleted uniformly} in the LB or, that outside of the LB the gas phase deuterium abundance may have been {\\it enhanced} along some LOS by the addition of nearly primordial gas which has recently fallen into the disk of the Galaxy in the form of cloudlets which take some time to mix with the pre-existing gas in the ISM.  Does \\ydlb~= 1.5 or \\ydmax~= 2.0 provide the best estimate of the lower bound to the ISM D abundance?  If deuterium is depleted onto dust, why is there not a strong correlation between deuterium abundance and iron depletion\\footnote{As pointed out by the Referee, shock strength may play a role in accounting for the scatter observed in the correlation between the gas phase D and Fe abundances.  If deuterium is loosely bound to the grain mantle while iron is locked into the core of the dust grain, deuterium would be more easily returned to the gas than iron when grains are exposed to shocks of modest strength, while iron might be removed from dust grains only by stronger shocks. The scatter in the correlation between the gas phase D and Fe abundances may be an indicator of shock strength.}  \\citep{linsky} and, which refractory element is then best to use as proxy for determining deuterium depletion onto dust?  These questions cannot be answered by the analysis presented here.  In a companion paper \\citep{sp}, abundances of refractory elements are used in concert with the deuterium abundances in an attempt to resolve this question.  \\subsection*"
    },
    "0910/0910.3770_arXiv.txt": {
        "abstract": "{ Baryonic Acoustic Oscillations (BAO) and their effects on the matter power spectrum can be studied by using the Lyman-$\\alpha$ absorption signature of the matter density field along quasar (QSO) lines of sight. A measurement sufficiently accurate to provide useful cosmological constraints requires the observation of $\\sim10^5$ quasars in the redshift range $2.2<z<3.5$ over $\\sim8000{\\rm deg^2}$. Such a survey is planned by the Baryon Oscillation Spectroscopic Survey (BOSS) project of the Sloan Digital Sky Survey (SDSS-III). } {One of the challenges for this project is to build from five-band imaging data a list of targets that contains the largest number of quasars in the required redshift range. In practice, one needs a stellar rejection of more than two orders of magnitude with a selection efficiency for quasars better than 50\\% up to magnitudes as large as $g\\sim 22$. Standard methods to identify quasars using colors work well for brighter quasars in the range 0.3~$<$~$z$~$<$~2.2 and $g<21$ but it is necessary to develop new methods for higher redshifts and magnitudes. } {To obtain an appropriate target list and estimate quasar redshifts, we have developed an Artificial Neural Networks (NN) with a multilayer perceptron architecture. The input variables are photometric measurements, i.e. the object magnitudes and their errors in the five bands ($ugriz$) of the SDSS photometry. The NN developed for target selection  provides a continuous output variable between 0 for non-quasar point-like objects to 1 for quasars. A second  NN  estimates the QSO  redshift $z$ using the photometric information.} {For target selection, we achieve a non-quasar point-like object rejection of 99.6\\% and 98.5\\% for a quasar efficiency of, respectively, 50\\% and 85\\%. The photometric redshift precision is of the order of 0.1 over the region relevant for BAO studies. These  statistical methods, developed in the context of the BOSS project, can easily be extended to any quasar selection and/or determination of their photometric redshift.  } {} ",
        "introduction": "Since the first quasar was discovered \\citep{schmidt63}, methods have been developed to differentiate these rare objects from other astronomical sources in the sky. In the standard methods, it is assumed that QSOs have point-like morphology. They are then separated from the much more numerous stars by their photometric colors. The UVX selection, e.g. \\citep{croom01}, can be largely complete ($>$90\\%) for QSOs with 0.3~$<$~$z$~$<$~2.2 but this completeness drops at higher redshift. The selection purity was brought up to 97\\% for $g<21$ using Kernel Density Estimation  techniques applied to SDSS colors \\citep{kde04} and extended to the infrared by \\citet{richards09} implying that spectroscopy is not needed to confirm the corresponding statistical sample of quasars at high galactic latitudes. This led to the definition of a one-million-QSO catalog \\citep{kde09} down to $i=21.3$ from the photometry of SDSS Data Release 6 \\citep{adelman08}.   Extending quasar selection methods to higher redshifts and magnitudes presents several difficulties. For example, at fainter magnitudes, galaxies start to contaminate ``point-like'' photometric catalogs both because of increasing photometric errors and because of non-negligible contributions of AGN's in certain bands. Nevertheless, such an extension is very desirable, not only to study the AGN population but also to use the quasars to study the foreground absorbers. In particular studies of spatial correlations in the IGM from the Lyman-$\\alpha$ forest and/or metal absorption lines  are in need of higher target density at high redshift \\citep{petitjean97,nusser99,pichon01,caucci08}.   More recently, it was realized that the Baryonic Acoustic Oscillations (BAO) could be detected in the Lyman-$\\alpha$ forest. BAO in the pre-recombination Universe imprint features in the matter power spectrum that have led to important constraints on the cosmological parameters. So far, BAO effects have been seen using galaxies of redshift $z<0.4$ to sample the matter  density \\citep{eisenstein,cole,percival}. The Baryon Oscillation Spectroscopic Survey (BOSS) \\citep{boss} of the Sloan Digital Sky Survey (SDSS-III) \\citep{sdss3} proposes to extend these studies using galaxies of higher redshifts, $z<0.9$. The BOSS project will also study BAO effects in the range $2.2<z<3.5$  using Lyman-$\\alpha$ absorption towards high redshift quasars (QSOs) to sample the matter density as proposed by \\citet{McDonald}.    \\begin{figure*}[htb] \\centering \\includegraphics[width=14.0cm]{Colors.eps} \\caption{2D distributions of colors ($u-g$, $g-r$, $r-i$, $i-z$ and  $g-i$)  for objects classified as PLO in  SDSS photometric catalog (blue lines for contours) and for  objects spectroscopically classified as QSO (red solid lines for contours). The PSF magnitudes ($ugriz$) have been corrected for Galactic extinction according to the model of \\citet{schlegel98}. } \\label{fig:Colors} \\end{figure*}  \\begin{figure*}[htb] \\centering \\includegraphics[width=14.0cm]{compVar.eps} \\caption{Distributions of the discriminating variables used as input in the NN for objects classified as PLO in  SDSS photometric catalog (blue dotted  histogram) and for  objects spectroscopically classified as QSO (red slashed histogram):  {\\bf a)} Distribution of the PSF $g$ magnitude, {\\bf b), c), d) e) and f)} Distributions of, respectively, $\\sigma(u)$, $\\sigma(g)$, $\\sigma(r)$, $\\sigma(i)$ and $\\sigma(z)$, the errors on the corresponding PSF magnitudes. } \\label{fig:compVar} \\end{figure*}   The power spectrum has already been measured at $z\\sim2.5$ via the 1-dimensional matter power spectrum derived from quasar spectra \\citep{croft}. The observation of BAO effects will require a full 3-dimensional sampling of the matter density, requiring a much higher number of quasars than previously available. BOSS aims to study around 100,000 QSOs over 8,000 square degrees. The requirement that the Lyman-$\\alpha$ absorption fall in the range of the BOSS spectrograph requires that the quasars be in the redshift range $2.2<z<3.5$.  The quasars to be targeted  must be chosen using only available photometric information, mostly from the SDSS-I point-source catalog. The target selection method must be able to reject the non-quasar point-like objects (PLOs; mainly stars) by more than two orders of magnitude with a selection efficiency of QSOs better than 50\\%. The BOSS project needs a high density of $z>2.2$ fainter QSOs ($\\sim 20$ QSOs per sq degree) and therefore requires the selection to be pushed up to $g\\sim 22$. We developed a new method to select quasars using more information than the standard color selection methods.   The classification of objects is a task that is generally performed by applying cuts on various distributions which distinguish signal objects from background objects. This approach is not  optimal  because all the information (the shapes of the variable distributions, the correlations between the variables) is not exploited and this leads to a loss in classification efficiency. Statistical methods based on multivariate analysis have been developed to tackle this kind of problem. For historical reasons these methods have been focused on linear problems which are easily tractable. In order to deal with nonlinearities, Artificial Neural Networks (NN) have been shown to be a powerful tool in the classification task (see for instance \\citet{bishop}).  By combining photometric measurements such as the magnitude values and their errors for the five bands ($ugriz$) of SDSS photometry, a NN approach will allow us both to select the QSO candidates and to predict their redshift. Similar methods such as Kernel Density Estimation (KDE) \\citep{kde04,kde09} already exist to select QSOs.  Our approach based on NN is an extension of these methods because we  will use more information (errors and absolute magnitude $g$ instead of only colors (difference between two magnitudes)). Moreover, we propose to  treat in parallel the determination of the redshift  with the same tool. This approach contrasts with the usual  methods to compute photometric redshift which deal with $\\chi^2$ minimization techniques  \\citep{photoz,weinstein04}.   \\section {QSO and Background Samples}    The quasar candidates should be selected among a photometric catalog of objects including real quasars and what we will call background objects. Here, both for the background and QSO samples, the photometric information comes from the SDSS-DR7 imaging database of point-like objects \\citep{dr7}, PLOs. We apply the same quality cuts on the photometry for the two samples and select objects with $g$ magnitude in the range $18 \\leq g  \\leq 22$. Note that in the following, magnitudes will be point spread function (PSF) magnitudes \\citep{lupton99} in the SDSS pseudo-AB magnitude system \\citep{oke83}.  \\subsection{Background Sample} \\label{sec:qsosample} For the background sample, we would ideally use an unbiased sample of spectroscopically confirmed SDSS point-like objects \\emph{that are not QSOs}. Unfortunately, we have no unbiased sample of such objects because spectroscopic targets were chosen in SDSS-I to favor particular types of objects. Fortunately, the number of QSOs among PLOs is sufficiently small that using all PLOs as background does not affect the NN's ability to identify QSOs. We have verified that this strategy works by using the synthetic PLO catalog of \\citet{fan}. We degraded the star sample by adding a few percent of QSOs in it. then, we retrained the NN and we  compared the NN trained with a pure star sample. We did not observe any significant worsening of the NN performances.    The background sample used in the following was drawn from the SDSS PLO sample. We used objects with galactic latitude $b$ around $45^\\circ$ to average the effect of Galactic extinction. In the future, we may consider the possibility of having a different NN for each stripe of constant galactic latitude. The final sample contains ~30,000 PLOs: half of them constituting the ``training\" sample, the other half  the ``control\" sample, as explained in the next section.   \\subsection{QSO Sample} For the QSO training sample, we use a list of 122,818 spectroscopically-confirmed quasars obtained from the 2QZ quasar catalog \\citep{croom04}, the SDSS-2dF LRG and QSO Survey (2SLAQ)  \\citep{croom09}, and the SDSS-DR7 spectroscopic database \\citep{dr7}. These quasars have redshifts in the range $  0.05 \\leq z  \\leq 5.0 $ and $g$ magnitudes  in the  range $18 \\leq g  \\leq 22$ (galactic extinction corrected). Since quasars will be observed over a limited blue wavelength range (down to about 3700~\\AA), we will target only quasars with $z>2.2$. Therefore, the sample of known quasars includes 33,918 QSOs with $z\\geq 1.8$: half of them constituting the effective ``training\" sample, the other half  the ``control\" sample. For the determination of the photometric redshift, we use a wider sample of 95,266 QSOs  with $z\\geq 1$.  In order to compare together QSOs with background objects from different regions of the sky, the QSO magnitudes have been corrected for Galactic extinction with  the model of \\citet{schlegel98}.  \\subsection{Discriminating variables}  The photometric information is extracted from the SDSS-DR7 imaging database \\citep{dr7}. The 10 elementary variables are the PSF magnitudes for the 5 SDSS bands ($ugriz$) and their errors. As explained in \\citet{kde09}, the most powerful  variables are the four usual colors ($u-g$,$g-r$,$r-i$,$i-z$) which combine the PSF magnitudes. Fig.~\\ref{fig:Colors} shows the 2D color-color distributions  for the QSO and PLO samples.  These plots give the impression that it is easy to disentangle the two classes of objects but one needs to keep in mind that the final goal is to obtain a 50\\% efficiency for QSOs with a non-quasar PLO efficiency of the order of $\\sim10^{-3}$. Therefore to improve the NN performances, we added the absolute magnitude $g$ and the five magnitude errors. Their distributions for the two classes are given on Fig.~\\ref{fig:compVar}. An improvement can be expected from the additional variables and also from the correlations between the variables. Indeed, for example, it is expected that errors be larger for compact galaxies compared to intrinsic point-like objects.   Note that the $g$ distribution for the QSOs is likely to be biased by the spectroscopic selection. This issue will be addressed in the future with the first observations of BOSS. Indeed the photometric selection of QSOs for these first observations is based on loose selection criteria and it should provide a ``less biased\" catalog of spectroscopically confirmed quasars, close to completeness up to $g=22$. ",
        "conclusions": "In this paper we have presented a new promising approach to select quasars from photometric catalogs and to estimate their redshift. We use an Neurone Network with a multilayer perceptron architecture. The input variables are photometric measurements, i.e. the magnitudes and their errors for the five bands ($ugriz$) of the SDSS photometry.  For the target selection, we achieve a PLO rejection factor of 99.6\\% and 98.5\\% for, respectively, a quasar efficiency of 50\\% and 85\\%.  The rms of the difference between the photometric redshift and the spectroscopic redshift is of the order of 0.15 over the region relevant for BAO studies. These new statistical methods developed in the context of the BOSS project can easily be extended to any other analysis requiring  QSO selection and/or  determination of their photometric redshift."
    },
    "0910/0910.1557_arXiv.txt": {
        "abstract": "{} {We establish the mean metallicity from high-resolution spectroscopy for the recently found dwarf spheroidal galaxy Bo\\\"otes\\,I and test whether it is a common feature for ultra-faint dwarf spheroidal galaxies to show signs of inhomogeneous chemical evolution (e.g. as found in the Hercules dwarf spheroidal galaxy).} {We analyse high-resolution, moderate signal-to-noise spectra for seven red giant stars in the Bo\\\"otes\\,I dSph galaxy using standard abundance analysis techniques. In particular, we assume local thermodynamic equilibrium and employ spherical model atmospheres and codes that take the sphericity of the star into account when calculating the elemental abundances.} {We confirm previous determinations of the mean metallicity of the Bo\\\"otes\\,I dwarf spheroidal galaxy to be --2.3 dex. Whilst five stars are clustered around this metallicity, one is significantly more metal-poor, at --2.9 dex, and one is more metal-rich at, --1.9 dex. Additionally, we find that one of the stars, Boo-127, shows an atypically high [Mg/Ca] ratio, indicative of  stochastic enrichment processes within the dSph galaxy. Similar results have previously only been found in the Hercules and Draco dSph galaxies and appear, so far, to be unique to this type of  galaxy. } {} ",
        "introduction": "Until recently, the number of dwarf spheroidal (dSph) galaxies around the Milky Way was small compared to expectations from $\\Lambda$CDM \\citep{moore1999}. However, in the past few years several new systems have been found through systematic searches \\citep[e.g.][]{belokurov.boo,belokurovcats}.  In general, dSph galaxies are some of the most tenuous stellar systems that we know of.  This is especially true for the newly found dSph galaxies which have very low stellar luminosities \\citep[see e.g., ][]{martin08}. The new dSphs are ultra-faint and show low metallicities as indicated by low-resolution spectroscopy \\citep{kirbyletter,kochbierman}. \\citet{koch2008Her} found unusual abundance patterns in two red giant stars (RGB) in the ultra-faint Hercules dSph galaxy. Because of the low baryonic mass for these system it has been speculated that the elemental abundances in the stars in these systems might show us the results of individual supernova events \\citep{koch2008Her}.  The recently found dSph galaxy in Bo\\\"otes \\citep[Bo\\\"otes\\,I, ][]{belokurov.boo} provides an excellent opportunity to test whether or not unusual elemental abundance ratios are a common feature of ultra-faint dSph galaxies, thanks to its low baryonic mass, \\citet{belokurov.boo} estimate, based on a colour magnitude diagram, that the Bo\\\"otes\\,I dSph galaxy is a purely old and metal-poor system.  Low-resolution spectroscopic data confirm this \\citep{martin07,norris2008} at find $<$[Fe/H]$>$=--2.5.  With $M_{\\rm V} \\sim -5.8$ this dSph galaxy is one of the least luminous galaxies known \\citep{belokurov.boo}. \\citet{fellhauer2008}, modelled the system and find that if this galaxy ever had a dark matter halo, it must still have it.  This implies that, since the dark matter provides a deep potential well, the stars that originally formed in the dSph galaxy are still there and, moreover, the depth of the well should have helped retain the ejecta from core-collapse supernova. For Hercules, \\citet{koch2008Her} conclude that about 10 supernova are needed to pollute the interstellar medium to the observed atypical abundance ratios. Given that Bo\\\"otes\\,I has an even lower baryonic mass than Hercules, we might expect to be able to see enrichment from individual supernovae in the elemental abundance trends (which would show up as large scatter in element ratios from star to star).  We have obtained high-resolution, moderate S/N spectra for seven RGB stars in the Bo\\\"otes\\,I dSph galaxy.  Here we report on the mean metallicity, the metallicity spread, and atypical abundance ratios similar to those found in the Hercules \\citep{koch2008Her} and Draco dSph galaxies \\citep{fulbright2004Draco}. Thus, Bo\\\"otes\\,I becomes the third system to show unexpected abundances ratios. ",
        "conclusions": ""
    },
    "0910/0910.4975_arXiv.txt": {
        "abstract": "We report on the method developed by Zibetti, Charlot \\& Rix (2009) to construct resolved stellar mass maps of galaxies from optical and NIR imaging. Accurate pixel-by-pixel colour information (specifically $g-i$ and $i-H$) is converted into stellar mass-to-light ratios with typical accuracy of 30\\%, based on median likelihoods derived from a Monte Carlo library of 50,000 stellar population synthesis models that include dust and updated TP-AGB phase prescriptions. Hence, surface mass densities are computed.  In a pilot study, we analyze 9 galaxies spanning a broad range of morphologies. Among the main results, we find that: i) galaxies appear much smoother in stellar mass maps than at any optical or NIR wavelength; ii) total stellar mass estimates based on unresolved photometry are biased low with respect to the integral of resolved stellar mass maps, by up to 40\\%, due to dust obscured regions being under-represented in global colours; iii) within a galaxy, {\\it on local scales} colours correlate with surface stellar mass density; iv) the slope and tightness of this correlation reflect/depend on the morphology of the galaxy. ",
        "introduction": "The key role of stellar mass to determine (or predict) the physical properties of present day galaxies has been established by a number of works since the last decade \\citep[e.g.][]{gavazzi+96,scodeggio+02,kauffmann+03b}. Indication has been provided that stellar mass density might be even more fundamental \\citep{bell_dejong00,kauffmann+03b}. All these works, however, deal with {\\it global} estimates of the stellar mass-to-light ratio, which is assumed to be uniform throughout a galaxy, contrary to what we should expect based on well known spatial variations of stellar population and dust properties. Hence, resolving the distribution of stellar mass is crucial to properly measure total stellar mass and mass density and to start investigating the structure and dynamics of galaxies in an unbiased way. Having access to the mass distribution also allows to address questions like: Is there any relation between {\\it local} stellar mass density and the {\\it local} SED and physical properties {\\it within} a galaxy, similarly to the relations observed {\\it globally}? If so, what does it tell us about the internal mechanisms of galaxy evolution?  With these goals and questions in mind, we have developed a new method to build stellar mass maps of galaxies \\citep[][ZCR09 hereafter]{ZCR09}, which we review in this contribution (Sec. 2). We also present preliminary results on the colour-stellar mass density relation within galaxies for a small sample of different morphological types (Sec. 3). ",
        "conclusions": ""
    },
    "0910/0910.1885_arXiv.txt": {
        "abstract": "The five dimensional Brans-Dicke theory naturally provides two scalar fields by the Killing reduction mechanism. These two scalar fields could account for the accelerated expansion of the universe. We test this model and constrain its parameter by using the type Ia supernova (SN Ia) data. We find that the best fit value of the 5-dimensional Brans-Dicke coupling contant is $\\omega = -1.9$. This result is also consistent with other observations such as the baryon acoustic oscillation (BAO). ",
        "introduction": "The expansion of the Universe is shown to be accelerating by the observations of type Ia supernovae (SN Ia) \\cite{Perlmutter99, Riess98}. This is usually attributed to the contribution of an unknown component, dubbed dark energy, which has negative pressure and makes up about three quarters of the total cosmic density (for recent measurements, see e.g. Ref.~\\cite{Komatsu08,Gong08}). The simplest model for dark energy is the cosmological constant (CC), which is consistent with most of the observations today. However, there are two big problems for CC, i.e. the well-known ``fine tuning problem'' and the ``coincidence problem'' . As alternatives, and also to solve these two problems, many dynamical dark energy models with scalar field have been proposed, such as quintessence \\cite{Wetterich88,Ratra88,FJ97,CDS98,CLW98, Carroll98,ZMS99,ZMS991}, phantom \\cite{Caldwell02}, quintom\\cite{FWZ05}, K-essence \\cite{Chiba00, Armendariz00}, tachyon \\cite{Padmanabhan02, Bagla03,Abramo03,Agui04,Guo04,Copeland05} and so on. Nevertheless, in most cases the fundamental physical origin of these scalar fields remain unknown, but just added by hand.  In Ref. \\cite{Qiang05}, by generalizing the Brans-Dicke theory to five dimensions and exploring its effect on the 4-dimensional world, another interesting approach to explain the cosmic accelerated expansion was proposed. Under the condition that the extra dimension is compact and sufficiently small, a spacelike Killing vector field $\\xi^a$ arises naturally, in which case the 5D Brans-Dicke theory can be reduced to a 4D theory, such that the 4-metric is coupled with two scalar fields $\\phi$ and $\\lambda$. Notes that here the scalar fields in four dimension stem naturally from a fundamental theory of gravity.  Considering the hypersurface-orthogonal property of $\\xi^a$, the line element in five dimension can take the form as \\begin{equation} ds^2=g_{\\mu\\nu}dx^{\\mu}dx^{\\nu}+\\lambda dx^5dx^5,\\label{metric} \\end{equation} thus the scalar $\\lambda$ also plays the role of a ``scale factor'' of the extra dimension. It was shown in Ref. \\cite{Qiang05} that these two scalar fields originated from the Killing reduction of the 5D Brans-Dicke theory may lead to the accelerated expansion of the universe. More detailed analysis is desirable to check if the theory can match the current observational data such as the SN Ia and the baryon acoustic oscillation.  In this paper, we compare the predictions of the cosmic expansion rate of this theory with the current cosmological observations, and constrain the 5D Brans-Dicke theory by means of the SN Ia data and the baryon acoustic oscillation (BAO) measurements. This work is based on the fact that the 4-dimensional gravitational constant $G$ varies extremely slowly with time \\cite{Turyshev04} in the current epoch. Thus we assume that $G$ is a constant at the ``low redshift'', so that the accretion of the white dwarf will not be affected, hence the luminosity and light curve of the observed SN Ia (redshifts range from 0 to 2) are not affected by the slow variations of $G$, and the SN Ia can still be used as the standard candle. Furthermore, if we assume that $G$ is almost a constant throughout the history of the Universe, we could also use the information from the large scale structure (e.g. BAO) to perform the constraints. However, we should note that $G=(\\phi\\lambda^{1/2}L)^{-1}$ \\cite{Qiang05}, where $L$ is the coordinate scale of the extra dimension. Although $G$ is almost constant in four dimension, $G^{(5)}\\sim\\phi^{-1}$ is not necessarily a constant and can still evolve with time. ",
        "conclusions": "By considering a hypersurface-orthogonal spacelike Killing vector field in the 5-dimensional spacetime, the 5D Brans-Dicke theory can be reduced to a 4D theory with the 4-metric coupled to two scalar fields. These two fields could naturally lead to the accelerated expansion of the Universe.  We study the evolution of the two fields and compare the expansion rate with SN Ia observations.  The two scalar field would make the Universe evolve as if ``matter-dominate'' or ``dark energy-dominate'' when $\\omega$ is greater or less than -2. We find that the model is in best agreement with the supernovae data when the 5-dimensional coupling constant $\\omega =-1.9 \\approx -2$, which happens to be also the value required to satisfy the solar system experiments. Furthermore, for this best fit value, the best fit $\\Omega_{m_0}$ value is about 0.27, in good agreement with other independent measurements such as those derived from X-ray cluster observations. This work is based on the assumption that the 4D gravitational constant $G$ varies extremely slowly so that it can be regarded as a constant at \"low redshift\" where the SN Ia data are available. If we further assume $G$ does not change during the whole history of the Universe, then other cosmological observations such as BAO can also be used, we find that in this case the results are almost the same.  In conclusion, the 5-dimensional Brans-Dicke theory could naturally provide two scalar fields which may cause the accelerated expansion, the result is consistent with the SN Ia observation, hence it is a candidate to explain the accelerated expansion of the Universe."
    },
    "0910/0910.3552_arXiv.txt": {
        "abstract": "A solution to the ultra-relativistic strong explosion problem with a non-power law density gradient is delineated. We consider a blast wave expanding into a density profile falling off as a steep radial power-law with small, spherically symmetric, and log-periodic density perturbations. We find discretely self-similar solutions to the perturbation equations and compare them to numerical simulations. These results are then generalized to encompass small spherically symmetric perturbations with arbitrary profiles. ",
        "introduction": "The profusion of explosions occurring in the observable universe has led to a pointed interest in the dynamics of blast waves. The quantitative treatment of strong explosions began with the self-similar solutions found by Sedov, Von-Neumann and Taylor \\cite{Sedov} \\cite{VonNeumann} \\cite{Taylor} for the flow behind spherical Newtonian shocks propagating into a cold gas with a power-law density profile, $\\rho \\propto r^{-k}$. This solution is valid for moderately steep decay exponents, $k<3$. Subsequently, corresponding solutions were found in the ultra-relativistic regime by Blandford \\& McKee \\cite{BM} for $k<4$. In these solutions the Lorentz factor of the shock $\\Gamma$ scales as $\\Gamma^2 \\propto t^{-m}$, and $m$ is fixed by energy conservation arguments. For $k>4$ this procedure fails due to the energy in the Blandford-McKee solution diverging, and a different argument must be used if a self-similar solution is to be found. It turns out that there exists a $k_g>4$ such that for $k>k_g$ just such an argument exists \\cite{Best}. The solutions in this regime are called type-II solutions \\cite{Zeldovich}, and what sets them apart from the solutions with $k<4$, known as type-I solutions, is the rapid acceleration of the shock and the fluid behind it that causes the formation of a sonic point between the shock and the center of the explosion. This point (actually a spherical surface) marks the boundary of an inner region that becomes causally disconnected from the shock. This allows a self-similar flow behind the shock to coexist with a non self-similar flow further away from it, thus maintaining a finite amount of total energy. For the ultra-relativistic case, $k_g=5-\\sqrt{3/4} \\approx 4.13$. The sonic point appears as a singularity in the hydrodynamic equations, and the requirement of regularity in traversing this singularity supplies the necessary condition to fix the  value of $m$.  In this paper we focus on Type-II ultra-relativistic solutions with $k>k_g$ \\cite{Best}, and use these as the basis for a perturbative analysis where we disturb the external density profile. The basic method for doing so was developed in a previous paper \\cite{Oren} for the Newtonian case, and we adapt it here for the ultra-relativistic case. We introduce a special family of perturbations to the ambient medium surrounding the explosion, so that we are able to reduce the hydrodynamic equations to a set of ordinary differential equations, and then find the flow behind the shock in the presence of perturbations. In section \\ref{sec:unpert} we briefly describe the unperturbed solutions, while in section \\ref{sec:RI} we shed light on some general properties of these solutions. In section \\ref{sec:pert} we write down the equations for the perturbations and solve them, and in section \\ref{sec:fourier} we generalize the results of the previous section by using a spectral decomposition of an arbitrary perturbation profile. Finally in section \\ref{sec:discuss} we summarize and discuss our results. ",
        "conclusions": "\\label{sec:discuss} We have laid out a method of solving the hydrodynamic equations for a strong explosion in the presence of spherically symmetric perturbations to the ambient density. At first we study a special group of perturbations with log-sinusoidal radial dependence, and discover an analytic solution. The requirement of spherical symmetry is not easily relaxed, because relativistic effects make it difficult to find self-similar solutions with a non trivial angular dependence. Another limitation is the linearity of the perturbation analysis that  limits the validity of the solutions to small amplitudes. The smallness required may be seen in figure \\ref{fig:sigmas} where numerical solutions with different amplitudes are superimposed against the analytical solution. It can be seen that our theory gives reasonably accurate results for amplitudes up to about $0.1$. The nonlinearity of waves with higher amplitudes is expressed through shock formation before their crests, as can be seen in the red line in figure \\ref{fig:sigmas}.  On the other hand we take advantage of linearity to generalize our results using a Fourier-like method, decomposing an arbitrary perturbation to simple modes for which we can solve the equations analytically. In this way we can treat interesting scenarios like a sudden rise or drop in density, which might e.g. be encountered in a stellar wind due to interaction with the interstellar medium. Another possible application is the emergence of a shock from the edge of a star \\cite{Pan}, where the drop in density accelerates shocks to relativistic velocities. The effect of perturbations in the star's envelope can be treated with the method presented here, requiring only an adaptation of the unperturbed solution.  Acknowledgments: The authors wish to thank Prof. Tsvi Piran for fruitful discussions. This research was partially supported by a NASA grant, IRG grant and a Packard Fellowship."
    },
    "0910/0910.2138_arXiv.txt": {
        "abstract": "\\label{abstract}\\label{firstpage} The nature of the progenitors of Type Ia supernovae (SNe Ia) is still unclear. In this paper, by considering the effect of the instability of accretion disk on the evolution of white dwarf (WD) binaries, we performed binary evolution calculations for about 2400 close WD binaries, in which a carbon--oxygen WD accretes material from a main-sequence star or a slightly evolved subgiant star (WD + MS channel), or a red-giant star (WD + RG channel) to increase its mass to the Chandrasekhar (Ch) mass limit. According to these calculations, we mapped out the initial parameters for SNe Ia in the orbital period--secondary mass ($\\log P^{\\rm i}-M^{\\rm i}_2$) plane for various WD masses for these two channels, respectively. We confirm that WDs in the WD + MS channel with a mass as low as $0.61\\,M_\\odot$ can accrete efficiently and reach the Ch limit, while the lowest WD mass for the WD + RG channel is $1.0\\,\\rm M_\\odot$. We have implemented these results in a binary population synthesis study to obtain the SN Ia birthrates and the evolution of SN Ia birthrates with time for both a constant star formation rate and a single starburst. We find that the Galactic SN Ia birthrate from the WD + MS channel is $\\sim$$1.8\\times 10^{-3}\\ {\\rm yr}^{-1}$ according to our standard model, which is higher than previous results. However, similar to previous studies, the birthrate from the WD + RG channel is still low ($\\sim$$3\\times 10^{-5}\\ {\\rm yr}^{-1}$). We also find that about one third of SNe Ia from the WD + MS channel and all SNe Ia from the WD + RG channel can contribute to the old populations ($\\ga$1\\,Gyr) of SN Ia progenitors. ",
        "introduction": "Type Ia supernovae (SNe Ia) play an important role in astrophysics, especially in cosmology. They appear to be good cosmological distance indicators and have been applied successfully to the task of determining cosmological parameters (e.g. $\\Omega$ and $\\Lambda$: Riess et al. 1998; Perlmutter et al. 1999). They are also a key part of our understanding of galactic chemical evolution owing to the main contribution of iron to their host galaxies (e.g. Greggio \\& Renzini 1983; Matteucci \\& Greggio 1986). Despite their importance, several key issues related to the nature of their progenitors and the physics of the explosion mechanisms are still not well understood (Hillebrandt \\& Niemeyer 2000; R\\\"{o}pke \\& Hillebrandt 2005; Podsiadlowski et al. 2008; Wang et al. 2008a), and no SN Ia progenitor system has been conclusively identified pre-explosion.  There is a broad agreement that SNe Ia are thermonuclear explosions of carbon--oxygen white dwarfs (CO WDs) in binaries (see the review of Nomoto, Iwamoto \\& Kishimoto 1997). Over the last few decades, two competing progenitor models of SNe Ia were discussed frequently, i.e. the double-degenerate (DD) and single-degenerate (SD) models. Of these two models, the SD model (Whelan \\& Iben 1973; Nomoto, Thielemann \\& Yokoi 1984; Fedorova, Tutukov \\& Yungelson 2004; Han 2008) is widely accepted at present. It is suggested that the DD model, which involves the merger of two CO WDs (Iben \\& Tutukov 1984; Webbink 1984; Han 1998), likely leads to an accretion-induced collapse rather than a SN Ia (Nomoto \\& Iben 1985; Saio \\& Nomoto 1985; Timmes, Woosley \\& Taam 1994). For the SD model, the companion is probably a main-sequence (MS) star or a slightly evolved subgiant star (WD + MS channel), or a red-giant star (WD + RG channel), or even a He star (WD + He star channel) (Hachisu, Kato \\& Nomoto 1996; Li $\\&$ van den Heuvel 1997; Hachisu et al. 1999a; Langer et al. 2000; Han $\\&$ Podsiadlowski 2004, 2006; Wang et al. 2009a,b). Although the SD model is currently a favorable progenitor model of SNe Ia, any single channel of the model cannot account for the birthrate inferred observationally. Note that, some recent observations have indirectly suggested that at least some SNe Ia can be produced by a variety of different progenitor systems (e.g. Hansen 2003; Ruiz-Lapuente et al. 2004; Patat et al. 2007; Voss \\& Nelemans 2008; Wang et al. 2008b; Justham et al. 2009).  At present, various progenitor models of SNe Ia can be examined by comparing the distribution of the delay time (between the star formation and SN Ia explosion) expected from a progenitor channel with that of observations (e.g. Chen $\\&$ Li 2007; Meng, Chen \\& Han 2009; L\\\"{u} et al. 2009; Ruiter, Belczynski \\& Fryer 2009). Mannucci, Della Valle \\& Panagia (2006) argued for the existence of two separate SN Ia populations, a `prompt' component with a delay time less than $\\sim$100\\,Myr, and a `delayed' component with a delay time $\\sim$3\\,Gyr. By investigating the star formation history (SFH) of 257 SN Ia host galaxies, Aubourg et al. (2008) found evidence of a short-lived population of SN Ia progenitors with lifetimes less than 180\\,Myr. Botticella et al. (2008) and Totani et al. (2008) analyzed host galaxies of SNe Ia and concluded that a substantial fraction of SNe Ia must have long delay times on the order of 2$-$3\\,Gyr. Moreover, Schawinski (2009) recently constrained the minimum delay time of 21 nearby SNe Ia by investigating their host galaxies (early-type galaxies). The study showed that these early-type host galaxies lack `prompt' SNe Ia with a delay time less than $\\sim$100\\,Myr and that $\\sim$70 per cent SNe Ia have minimum delay times of 275\\,Myr$-$1.25\\,Gyr, while at least 20 per cent SNe Ia have minimum delay times of at least 1\\,Gyr at 95 per cent confidence and two of these four SNe Ia are likely older than 2\\,Gyr.  For SNe Ia with the short delay times, Wang et al. (2009a) studied a WD + He star channel to produce SNe Ia, in which a CO WD accretes material from a He MS star or a He subgiant to increase its mass to the Chandrasekhar (Ch) mass. The study showed the parameter spaces for the progenitors of SNe Ia. By using a detailed binary population synthesis (BPS) approach, Wang et al. (2009b, WCMH09) found that the Galactic SN Ia birthrate from this channel is $\\sim$$0.3\\times 10^{-3}\\ {\\rm yr}^{-1}$, and that this channel can produce SNe Ia with short delay times ($\\sim$45$-$140\\,Myr). For SNe Ia with the long delay times ($\\ga$1\\,Gyr), this requires that the mass of the companion should be $\\la$$2\\,M_{\\odot}$.  Recently, Xu \\& Li (2009) emphasized that the mass-transfer through the Roche lobe overflow (RLOF) in the evolution of WD binaries may become unstable (at least during part of the mass-transfer lifetime), i.e. the mass-transfer rate is not equivalent to the mass-accretion rate onto the WD. This important feature has been ignored in nearly all of the previous theoretical works on SN Ia progenitors except for King, Rolfe \\& Schenker (2003) \\footnote{King, Rolfe \\& Schenker (2003) adopted a similar method in Li \\& Wang (1998) to produce SNe Ia with long period dwarf novae in a semi-analytic approach.} and Xu \\& Li (2009), who inferred that the mass-accretion rate onto the WD during dwarf nova outbursts can be sufficiently high to allow steady nuclear burning of the accreted matter and growth of the WD mass. In particular, the study of Xu \\& Li (2009) enlarges the region of SN Ia parameter spaces. However, they only give the SN Ia parameter spaces with WD initial masses of 0.8 and $1\\,M_{\\odot}$. More detailed work is obviously needed to investigate the influence of the accretion disk on the final results and to give SN Ia birthrates and delay times.   Including the effect of the instability of accretion disk on the evolution of WD binaries, the purpose of this paper is to study the WD + MS and WD + RG channels towards SNe Ia comprehensively and systematically, and then to determine the parameter spaces for SNe Ia, which can be used in BPS studies. In Section 2, we describe the numerical code for the binary evolution calculations and the grid of the binary models. The binary evolutionary results are shown in Section 3. We describe the BPS method in Section 4 and present the BPS results in section 5. Finally, a discussion is given in Section 6, and a summary in Section 7. ",
        "conclusions": "\\label{6. DISCUSSION} The regions (Fig. 6 in this paper) for producing SNe Ia depend on many uncertain input parameters, in particular for the duty cycle which is poorly known. The main uncertainties lie in the facts that it varies from one binary system to another and may evolve with the orbital periods and mass-transfer rates (e.g. Lasota 2001; Xu \\& Li 2009). This is the reason why we choose an intermediate value (0.01) rather than other extreme ones (e.g. 0.1 or $10^{-3}$). However, we also did some tests for a higher or lower value of the duty cycle. We find that the variation of the duty cycle value will influence the regions for producing SNe Ia, especially for the WD + RG channel (for a high value (0.1), the regions will be shifted to higher period; for a low value ($10^{-3}$), the regions will be shifted to lower period). For these two extreme duty cycle values, the SN Ia birthrate from the WD + RG channel will be lower due to the smaller regions for producing SNe Ia.  In our binary evolution calculations we have not considered the influence of rotation on the H-accreting WDs. The calculations of Yoon, Langer \\& Scheithauer (2004) showed that the He-shell burning is much less violent when rotation is taken into account. This may significantly increase the He-accretion efficiency ($\\eta _{\\rm He}$ in this paper). Meanwhile, the maximum stable mass of a rotating WD may be above the standard Ch mass (i.e. the super-Ch mass model: Uenishi, Nomoto \\& Hachisu 2003; Yoon \\& Langer 2005; Chen \\& Li 2009). However, we mainly focus on the standard Ch mass explosions of the accreting WDs in this work.  In our BPS studies we assume that all stars are in binaries and about 50\\,per cent of stellar systems have orbital periods less than 100\\,yr. In fact, this is known to be a simplification, and the binary fractions may depend on metallicity, environment, spectral type, etc. If we adopt 40\\,per cent of stellar systems have orbital periods below 100\\,yr by adjusting the parameters in Equation (10), we estimate that the Galactic SN Ia birthrate from the WD + MS channel will decrease to be $\\sim$$1.4\\times 10^{-3}\\ {\\rm yr}^{-1}$ according to our standard model (the birthrate from the WD + RG channel will decrease to be $\\sim$$2\\times 10^{-5}\\ {\\rm yr}^{-1}$). In addition, Umeda et al. (1999) concluded that the upper limit mass of CO cores born in binaries is about 1.07$\\,M_\\odot$. If this value is adopted as the upper limit of the CO WD, the SN Ia birthrate from the WD + MS channel will decrease to be $\\sim$$1.7\\times10^{-3}\\,{\\rm yr}^{-1}$ (the birthrate from the WD + RG channel will decrease to be $\\sim$$1\\times 10^{-5}\\ {\\rm yr}^{-1}$).  The Galactic SN Ia birthrate from the WD + RG channel is $\\sim$3$\\times 10^{-5}\\ {\\rm yr}^{-1}$ according to our standard model, which is low compared with observations, i.e. SNe Ia from this channel may be rare. However, further study on this channel is necessary, since this channel may explain some SNe Ia with long delay times. In addition, it is suggested that, RS Oph and T CrB, both recurrent novae are probable SN Ia progenitors and belong to the WD + RG channel (e.g. Belczy$\\acute{\\rm n}$ski \\& Mikolajewska 1998; Hachisu, Kato \\& Nomoto 1999b; Sokoloski et al. 2006; Hachisu, Kato \\& Luna 2007). Meanwhile, by detecting Na I absorption lines with low expansion velocities, Patat et al. (2007) suggested that the companion of the progenitor of SN 2006X may be an early RG star. Additionally, Voss \\& Nelemans (2008) studied the pre-explosion archival X-ray images at the position of the recent SN 2007on, and they considered that its progenitor may be a WD + RG system.     Employing Eggleton's stellar evolution code with the prescription of Hachisu et al. (1999a) for the mass-accretion of CO WDs, and including the effect of the instability of accretion disk on the evolution of WD binaries, we performed binary evolution calculations for about 2400 close WD binaries. The calculated results further confirm that the disk instability could substantially increase the mass-accumulation efficiency for accreting WDs, and cause the possible SNe Ia to occur in systems with $\\la$$2\\,M_{\\odot}$ donor stars (see also King, Rolfe \\& Schenker 2003; Xu \\& Li 2009). We find that the Galactic SN Ia birthrate from the WD + MS channel is $\\sim$$1.8\\times 10^{-3}\\ {\\rm yr}^{-1}$ according to our standard model, which is higher than previous results. However, similar to previous studies, the birthrate from the WD + RG channel is still low ($\\sim$$3\\times 10^{-5}\\ {\\rm yr}^{-1}$). We also find that about one third of SNe Ia from the WD + MS channel and all SNe Ia from the WD + RG channel can contribute to the old populations ($\\ga$1\\,Gyr) of SN Ia progenitors. The companion stars of SNe Ia with long delay times in this work would survive in the SN explosion and show distinguishing properties. In future investigations, we will explore the properties of the companion stars after SN explosion, which could be verified by future observations."
    },
    "0910/0910.5718_arXiv.txt": {
        "abstract": "We present a study of galaxy mergers and the influence of environment in the Abell 901/902 supercluster at $z\\sim$~0.165, based on 893 bright ($R_{\\rm Vega} \\le$~24) intermediate mass ($M_{*} \\geq 10^{9} M_{\\sun}$) galaxies.  We use $HST$ ACS F606W data from the Space Telescope A901/902 Galaxy Evolution Survey (STAGES), COMBO-17, $Spitzer$ 24$\\micron$, and $XMM$-$Newton$ X-ray data.  Our analysis utilizes both a physically driven visual classification system, and quantitative CAS parameters to identify systems which show evidence of a recent or ongoing merger of mass ratio $ >$~1/10 (i.e., major and minor mergers).  Our results are: (1)~After visual classification and minimizing the contamination from false projection pairs, we find that the merger fraction $f_{\\rm merge}$ is 0.023$\\pm$0.007.  The estimated fractions of likely major mergers, likely minor mergers, and ambiguous cases are 0.01$\\pm$0.004, 0.006$\\pm$0.003, and 0.007$\\pm$0.003, respectively.  (2)~All the mergers lie outside the cluster core of radius $R <$~0.25 Mpc: the lack of mergers in the core is likely due to the large galaxy velocity dispersion in the core. The mergers, instead, populate the region (0.25 Mpc $< R\\leq$~2 Mpc) between the core and the cluster outskirt. In this region, the estimated frequency of mergers is similar to those seen at typical group overdensities in N-body simulations of accreting groups in the A901/902 clusters. This suggests ongoing growth of the clusters via accretion of group and field galaxies.  (3)~We compare our observed merger fraction with those reported in other clusters and groups out to $z\\sim$~0.4. Existing data points on the merger fraction for $L \\leq L^{*}$ galaxies in clusters allow for a wide spectrum of scenarios, ranging from no evolution to evolution by a factor of $\\sim$5 over $z\\sim$~0.17 to 0.4.  (4)~In A901/902, the fraction of mergers, which lie on the blue cloud is 80$\\%\\pm 18\\%$ (16/20) versus 34$\\%\\pm7\\%$ or (294/866) for non-interacting galaxies, implying that interacting galaxies are preferentially blue.  (5)~The average SFR, based on UV or a combination of UV+IR data, is enhanced by a factor of $\\sim$1.5 to 2 in mergers compared to non-interacting galaxies.  However, mergers in the A901/902 clusters contribute only a small fraction (between 10\\% and 15\\%) of the total SFR density, while the rest of the SFR density comes from non-interacting galaxies. ",
        "introduction": "Understanding how galaxies evolve in various environments (field, groups, and clusters), and as a function of redshift is a key step toward developing a coherent picture of galaxy evolution.  Present-day cluster and field galaxies differ due to several factors, which are often grouped under the umbrella of `nature' versus `nurture'. First, in cold dark matter (CDM) cosmogonies, the first galaxies formed and evolved early in cluster cores, as the higher initial overdensities led to faster gravitational collapse and more rapid mergers of proto-galaxies (e.g., Cole et \\etal 2000; Steinmetz \\& Navarro 2002).  Second, in the context of the bottom-up CDM assembly paradigm, the outer parts of clusters and superclusters, grow at late times via mergers, smooth accretion, and discrete accretion of groups and field galaxies.  This idea is supported by observational studies (e.g., Zabludoff \\& Franx 1993; Abraham \\etal 1996a; Balogh, Navarro \\& Morris 2000), which suggest that clusters continuously grew by the accretion of groups.  Third, the dominant physical processes affecting galaxies differ in cluster and field environments due to the different galaxy number density, galaxy velocity dispersion, and intracluster medium (ICM).  Among these processes are close galaxy-galaxy interactions, such as strong tidal interactions and mergers (e.g., Barnes 1992; Moore \\etal 1998) and galaxy harassment (e.g., Moore \\etal 1996), which stems from the cumulative effect of weak interactions.  Furthermore, in clusters where the hot ICM makes up as much as 15\\% of the total mass, galaxy-ICM interactions, such as ram pressure stripping (Gunn \\& Gott 1972; Larson \\etal 1980; Quilis \\etal 2000; Balogh, Navarro \\& Morris 2000), can play an important role in removing the diffuse gas from galaxies. The tidal field of the cluster potential may also play a relevant role in the dynamical evolution of cluster galaxies (Gnedin 2003).   Systematic studies of the differences between cluster, group, and field galaxies at different redshifts are needed to shed light on the relative importance of these various processes in their respective environments.  Several differences have been observed between galaxies in the field and those in the rich cluster environment, but the physical drivers behind these variations are still under investigation.  At $z\\sim$~0, the relative percent of massive early type (E+S0) galaxies to spirals rises from (10\\%+10\\%:80\\%) in the field to (40\\%+50\\%:10\\%) in the cores of very rich clusters, leading to the so-called morphology density relation (Dressler 1980; Dressler \\etal 1997).  However, recent Sloan Digital Sky Survey (SDSS) studies suggest that masses and star formation (SF) histories of galaxies are more closely related to environmental physical processes rather than their structural properties (Blanton \\etal 2005).  The SF histories of galaxies depend on both luminosity (Cole et al. 2001) and environment (Diaferio \\etal 2001; Koopmann \\& Kenney 2004).  The fraction of blue galaxies in clusters appears to rise with redshift, an effect known as the Butcher Oemler effect (Butcher \\& Oemler 1978; Margoniner \\etal 2001; de Propis \\etal 2003).  There is also evidence that SF in bright ($M_{\\rm v} < $~-18) cluster galaxies is suppressed compared to field galaxies (e.g., Balogh \\etal 1998, 1999), for reasons that are not well understood.   Galaxy interactions and mergers have been proposed as a mechanism for the change in galaxy populations in clusters from that of the field (e.g., Toomre \\& Toomre 1972; Lavery \\& Henry 1988; Lavery Pierce, \\& McClure 1993).  There have been various studies on the properties of galaxies (e.g., Dressler 1980, Postman \\& Geller 1984, Giovanelli, Haynes, \\& Chincarini, 1986, Kennicutt 1983; Gavazzi \\& Jaffe 1985, Whitmore \\etal 1993) and of galaxy interactions and mergers (e.g., Lavery \\& Henry 1988; Lavery, Pierce, \\& Mclure 1992; Zepf 1993; Dressler \\etal 1994; Couch \\etal 1998; van Dokkum \\etal 1998, 1999; Tran \\etal 2005, 2008) in different environments.  Some of these studies suggest that galaxy interactions and mergers may play a role in morphological transformations of galaxies in clusters, but there have been few systematic studies of galaxy mergers and interactions in clusters, based on high resolution $HST$ images as well as $Spitzer$ 24$\\micron$, and X-ray images.  In this paper we present a study of the frequency, distribution, color, and SF properties of galaxy mergers in the A901/902 supercluster at $z\\sim$0.165. We use $HST$ ACS F606W data taken as part of the Space Telescope A901/902 Galaxy Evolution Survey (STAGES; Gray \\etal 2009), along with ground-based COMBO-17 imaging data (Wolf \\etal 2004), $Spitzer$ 24$\\micron$ data (Bell \\etal 2005, 2007), $XMM$-$Newton$ X-ray data (Gilmour \\etal 2007), and dark matter (DM) mass measurements from weak lensing (Heymans \\etal 2008). With a resolution of $0.1\\arcsec$ or $\\sim$280 pc at $z=0.165$, the $HST$ images allow for the identification of merger signatures such as double nuclei, arcs, shells, tails, tidal debris, and accreting satellites.  The COMBO-17 survey (Wolf \\etal 2004) provides accurate spectrophotometric redshifts down to $R_{\\rm Vega}$ of 24 and stellar masses (Borch \\etal 2006).  The $Spitzer$ 24$\\micron$ data (Bell \\etal 2005, 2007) probe the obscured SF, while X-ray maps (Gilmour \\etal 2007; Gray \\etal, 2010 in preparation) provide information of how the ICM density changes throughout the cluster.    We present the data and sample selection in $\\S$Section~\\ref{sdatas}.  In $\\S$Section~\\ref{smvc} and $\\S$Section~\\ref{smcas1}, we describe the two different methods that we use to identify galaxy mergers: a physically motivated classification system that uses visual morphologies, stellar masses, and spectrophotometric redshifts, and a method based on CAS asymmetry $A$ and clumpiness $S$ parameters (Conselice \\etal 2000). In $\\S$Section~\\ref{sinte1} and~\\ref{scas1}, we explore the frequency of galaxy mergers in A901/902 based on these two methods and present one of the first systematic comparisons to date between CAS-based and visual classification results in clusters.  We set a lower limit on the fraction of major mergers (those with mass ratio $M_{1}/M_{2} \\ge 1/4$).  In $\\S$Section~\\ref{sinte2}, we examine the distribution of mergers in the A901/902 supercluster as a function of clustocentric radius, galaxy number density, local galaxy surface density, relative ICM density, and local DM mass surface density.  In $\\S$Section~\\ref{sinte3}, we compare our results on the fraction and distribution of mergers to expectations based on analytical estimates and simulations of mergers in different environments.  In $\\S$Section~\\ref{scomp1}, we compare our results on galaxy mergers in the A901/902 supercluster to groups and clusters at different redshifts out to $z\\sim$~0.8.  We investigate the fraction of mergers and non-interacting galaxies on the blue cloud and red sample as a function of clustocentric radius in $\\S$Section~\\ref{scolor1}.  Finally in $\\S$Section~\\ref{ssfr1}, we compare the star formation rate (SFR) of mergers and non-interacting galaxies in the A901/902 clusters. The results of this work are summarized in $\\S$Section~\\ref{ssumm}.  In this paper, we assume a flat cosmology with $\\Omega_{\\rm m}$ = 1 $-\\Omega_{\\lambda}$ = 0.3 and H$_{0}$~=~70 km s$^{-1}$ Mpc$^{-1}$ throughout this paper. ",
        "conclusions": "\\label{sresul}  \\subsection{Merger fraction in A901/902 from visual classification}\\label{sinte1}  Next we discuss how to define and estimate the merger fraction.  Our goal is to estimate the fraction of $f_{\\rm merge}$ of systems with stellar mass above an appropriately chosen mass cut $M_{\\rm cut}$, which are involved in mergers of mass ratio $\\ge$~1/10.  The merger fraction $f_{\\rm merge}$ is computed as ($N_{\\rm merge}/N_{\\rm tot}$), where $N_{\\rm merge}$ is the number of major and minor mergers involving galaxies with $M_{*} \\ge M_{\\rm cut}$, and $N_{\\rm tot}$ is the total number of galaxies with $M_{*} \\ge M_{\\rm cut}$.   It is important to determine the minimum stellar masses of the major and minor mergers involving galaxies with $M_{*} \\ge M_{\\rm cut}$ and thereby assess for what value of $M_{\\rm cut}$ we can trace such mergers.  Major merger pairs (defined as having mass ratio 1/4 $< M_{1}/M_{2} \\le$ 1) of mass ratio 1:1 to 1:3 will have minimum stellar masses ranging from $2 \\times M_{\\rm cut}$ to $4 \\times M_{\\rm cut}$, in cases where galaxies of mass $M_{*} \\ge M_{\\rm cut}$ merge with systems at least as massive as themselves.  However, if galaxies of mass $M_{*} \\ge M_{\\rm cut}$ merge with lower mass systems, the minimum mass of 1:1 to 1:3 mergers will range from $2 \\times M_{\\rm cut}$ to ($4 \\times M_{\\rm cut}/3$).  Taken together, the above constraints imply that we are complete for major mergers involving galaxies with $M_{*} \\ge M_{\\rm cut}$, as long as we can trace {\\rm 1:3 major mergers with a minimum stellar mass of ($4 \\times M_{\\rm cut}/3$).  Repeating the exercise for minor mergers (defined as having 1/10 $< M_{1}/M_{2} \\le$ 1/4), it follows that the minimum stellar mass for 1:4 to 1:9 mergers ranges from $5 \\times M_{\\rm cut}$ to $9 \\times M_{\\rm cut}$ or from ($5 \\times M_{\\rm cut}/4$) to ($10 \\times M_{\\rm cut}/9$), depending on whether galaxies of mass $M_{*} \\ge M_{\\rm cut}$ merge with systems of higher mass or lower mass.  Thus, we are complete for minor mergers involving galaxies with $M_{*} \\ge M_{\\rm cut}$, as long as we can trace {\\rm 1:9 minor mergers with a minimum stellar mass of ($10 \\times M_{\\rm cut}/9$).   For what value of $M_{\\rm cut}$ are these conditions satisfied by the mergers of Types 1, 2a, and 2b, which we visually identified in the final sample ($\\S$Section~\\ref{smvc})?  Since our final sample is complete for $M_{*} \\ge 10^{9} M_{\\sun}$, it follows that we can identify all mergers of type 1 and 2a with $M_{*} \\ge 10^{9} M_{\\sun}$ from this sample.  If we impose this value to the afore-defined criteria of ($10 \\times M_{\\rm cut}/9$) and ($4 \\times M_{\\rm cut}/3$), it implies that a mass cut $ M_{\\rm cut} \\ge 0.9 \\times 10^{9} M_{\\sun}$ would allow us to be complete for major and minor mergers of type 1 and 2a. However, the situation is different for the mergers of type 2b (close pairs resolved into two galaxies by COMBO-17) because the individual galaxies making up the pair are complete only for $M _{*} \\geq10^{9}$ $M_{\\sun}$. As a result, we can only completely trace 1:3 major mergers of type 2b with total mass $\\ge 4.0 \\times 10^{9} M_{\\sun}$. If we impose this value to the afore-defined criterion of ($4 \\times M_{\\rm cut}/3$)}, it implies that a mass cut of $ M_{\\rm cut} \\ge 3.0 \\times 10^{9} M_{\\sun}$ is needed to ensure completeness for major mergers of type 2b.  Similarly, we can only completely trace 1:9 minor mergers of type 2b with total mass $ \\ge 9.0 \\times 10^{9} M_{\\sun}$, which implies that a mass cut $M_{\\rm cut} \\ge 9.0 \\times 10^{9} M_{\\sun}$ is needed to completely trace all minor mergers of type 2b.  Using a mass cut $M_{\\rm cut} \\ge 9.0 \\times 10^{9} M_{\\sun}$ for computing the merger fraction $f_{\\rm merge}$ would allow us to be complete for major and minor mergers of types 1, 2a, and 2b.  However, it leads to very small number statistics and is not viable.  We therefore explore the value of $f_{\\rm merge}$ for the less severe cut of $M_{\\rm cut} \\ge 3.0 \\times 10^{9} M_{\\sun}$, as well as for $M_{\\rm cut} \\ge 10^{9} M_{\\sun}$.  The cut of $M_{\\rm cut} \\ge 3.0 \\times 10^{9} M_{\\sun}$ ensures that we are complete for major and minor mergers of type 1, 2a, but leaves us incomplete for mergers of type 2b.  It is encouraging that both cuts yield similar values for the merger fraction $f_{\\rm merge}$, as illustrated in Table~\\ref{tcut}: the merger fraction is 0.023$\\pm$0.007 and 0.021$\\pm$0.007, respectively, for $M_{\\rm cut} \\ge 10^{9} M_{\\sun}$ and $M_{\\rm cut} \\ge 3.0 \\times 10^{9} M_{\\sun}$.  Note that in computing the merger fraction, we only include the 20 distorted mergers listed in Table~\\ref{tmer}, and avoid the potential projection pairs of type 2a and 2b without signs of morphological distortions \\footnote{Even if we included all pairs of type 2a and 2b, the merger fraction would still have similar values (0.041$\\pm$0.01 and 0.033$\\pm$0.01) for both mass cuts (Table~\\ref{tcut}).}.   While mergers may have played an important role in the evolution of cluster galaxies at earlier times (e.g., $z>2$), hierarchical models (e.g., Gottloeber et al 2001; Khochfar \\& Burkert 2001) predict that the merger fraction in dense clusters falls more steeply at $z<1$ than the field merger fraction.  As a result, at $z<0.3$, the merger fraction for intermediate mass cluster galaxies is predicted to be quite low (typically below 5\\%).  The low merger fraction among intermediate mass ($M = 10^{9}$ to a few $\\times 10^{10} M_{\\sun}$) galaxies in the A901/902 clusters is consistent with the latter prediction.  How are the mergers distributed among major mergers, minor mergers, and ambiguous cases that could be either major or minor mergers?  The results are shown in Columns 8--10 of Table ~\\ref{tcut}, based on the classification listed in Column 8 of Table ~\\ref{tmer}.  The estimated fractions of likely major mergers, likely minor mergers, and ambiguous cases are 0.01$\\pm$0.004\\% (9/886), 0.006$\\pm$0.003\\% (5/886), and 0.007$\\pm$0.003\\% (6/886), respectively for $M_{\\rm cut} \\ge 10^{9} M_{\\sun}$.  For $M_{\\rm cut} \\ge 3.0 \\times 10^{9} M_{\\sun}$, the corresponding fractions are 0.013$\\pm$ 0.005\\% (8/609), 0.008$\\pm$0.004\\% (5/609), and 0.0\\% (0/609), respectively.   In the rest of this paper, we continue to work with a mass cut of $M_{\\rm cut} \\ge 10^{9} M_{\\sun}$, and the 20 distorted mergers (Table~\\ref{tmer}) applicable for this mass cut.  However, where relevant, we cite many of our results for both of the mass cuts ($M_{\\rm cut} \\ge 10^{9} M_{\\sun}$ and $M_{\\rm cut} \\ge 3.0 \\times 10^{9} M_{\\sun}$) so that we can gauge the potential effect of incompleteness in tracing major and minor mergers.     \\subsection{Frequency of mergers in A901/902 from CAS}\\label{scas1} The results of running CAS on the final classified sample of 886 intermediate-mass ($M_{*} \\geq 10^{9}$ $M_{\\sun}$) systems are shown in Table~\\ref{tcas} and Figures~\\ref{fcas1} and~\\ref{fcas2}.  In Figure~\\ref{fcas1}, the 20 mergers of type 1, 2a, and 2b are plotted in different symbols.  For Type 2b merger pairs, we plot the highest value of CAS $A$ found between the two galaxies in each system.  Using the CAS merger criterion ($A>0.35$ and $A>S$) to identify mergers yields a merger fraction ($f_{\\rm CAS}$) in this sample of 18/886 or $0.02\\pm 0.006$.  For the more conservative sample with a mass cut $M_{\\rm cut} \\ge 3 \\times 10^{9} M_{\\sun}$, $f_{\\rm CAS}$ is 7/609 or $0.011\\pm0.005$.  When citing the error on $f_{\\rm CAS}$, we take the largest of either the Poisson error or the systematic error ($\\sigma_{\\rm CAS}$) in $f_{\\rm CAS}$ due to the systematic errors in CAS $A$ and $S$.  The systematic error, $\\sigma_{\\rm CAS}$, is calculated by taking the upper and lower bounds of $f_{\\rm CAS}$ based on ($ A\\pm$ error in $A$) and ($S \\pm$ error in $S$). Specifically, these limits on $f_{\\rm CAS}$ are found by using the criteria of (($A \\pm$ error in $A$) $<$ ($S \\pm$ error in $S$)) and (($A \\pm$ error in $A$) $>$ 0.35).  At first sight, the CAS-based merger fractions $f_{\\rm CAS}$ for the two mass cuts ($0.020\\pm0.006$ and $0.011\\pm0.005$ for $M_{\\rm cut} \\ge 10^{9} M_{\\sun}$ and $M_{\\rm cut} \\ge 3\\times 10^{9} M_{\\sun}$, respectively) are not widely different from the merger fraction $f_{\\rm merge}$ based on visual classification (0.023$\\pm$0.007 and 0.021$\\pm$0.007 for $M_{\\rm cut} \\ge 10^{9} M_{\\sun}$ and $M_{\\rm cut} \\ge 3\\times 10^{9} M_{\\sun}$, respectively; $\\S$Section~\\ref{sinte1}). However, the comparison of $f_{\\rm CAS}$ and $f_{\\rm merge}$ does not tell the whole story because the nature of the systems picked by the two methods can be quite different.  The visual classes of the 18 systems, which satisfy the CAS criterion and are considered as mergers in the CAS system, are shown in Table~\\ref{tcas}.  It turns out 7/18 ($39\\pm14\\%$) of these ``CAS mergers'' are visually classified as non-interacting systems.  The results are illustrated in Table~\\ref{tcas} and Figure~\\ref{fcas1}.  Figure~\\ref{fcas2} shows examples of these ``contaminants'': they tend to be dusty highly inclined systems and systems with low level asymmetries that seem to be caused by SF.  It is also useful to ask what fraction of the 20 systems visually-identified as mergers satisfy the CAS criterion.  We find that for $M_{\\rm cut} \\ge 10^{9} M_{\\sun}$, the CAS criterion only captures 11 of the 20 ($55\\pm16\\%$) of the visually-classified mergers. These results are illustrated in Figure~\\ref{fcas1}. Figure~\\ref{fcas2} shows examples of merging galaxies missed by CAS. The missed mergers include galaxies with fainter outer tidal features, double nuclei where CAS puts the center between the nuclei, and pairs of very close connected galaxies.  It is also interesting to ask what percentage of the different Types of mergers (Types 1, 2a, and 2b) does the CAS merger criterion recover.  Figure~\\ref{fcas1} also shows the mergers divided up by merger types 1, 2a and 2b.  Of the 13 mergers of type 1, CAS recovers 9/13 or 69\\%$\\pm$18\\%. Of the 3 mergers of type 2a, CAS recovers 1/3 or 33\\%$\\pm$28\\%. Of the 4 mergers of type 2b, CAS recovers 1/4 or 25\\%$\\pm$22\\%.  Since the CAS criterion is widely used to pick major mergers, it is also interesting to explore how well this criterion picks up the systems that we visually classified as major mergers (Table~\\ref{tmer}).  We find that the CAS merger criterion picks up 6/9 (67\\%$\\pm20\\%$) of the systems classified as major mergers.  It is also interesting to note that the CAS criterion recovers 2/5 (40\\%$\\pm23\\%$) and 4/6 (67\\%$\\pm23\\%$) of the systems classified respectively as minor mergers and ambiguous mergers.     \\subsection{Distribution of  mergers}\\label{sinte2}   In order to define different regions of the A901a, A901b, and A902 clusters, we computed the projected number density $n$ (Figure~\\ref{fnumd}) for intermediate-mass ($M_{*} \\geq 10^{9} M_{\\sun}$) galaxies as a function of clustocentric radius $R$ by assuming a spherical distribution.  Note here that each galaxy is assigned to the cluster closest to it, and $R$ is measured from the center of that cluster.  We consider the cluster core to be at $R\\le$~0.25 Mpc, as this is the region where the projected number density $n$ rises very steeply (Figure~\\ref{fnumd}).  The cluster virial radii are taken to be $\\sim$1.2 Mpc, based on estimates from the DM maps derived from gravitational lensing by Heymans \\etal (2008).  Throughout this paper, we refer to the region at 0.25 Mpc~$<R \\le$~1.2 Mpc, between the cluster core and the cluster virial radius, as the `outer region of the cluster'. The region outside the virial radius (1.2 Mpc~$< R \\le$~2.0 Mpc) is referred to as the `outskirt region of the cluster'. The core, outer region, and outskirt region are labeled on Figure~\\ref{fnumd}.   As shown in Table~\\ref{tcompa}, the A901 clusters have central galaxy number densities ($n$) in the core region of 1000--1600 galaxies Mpc$^{-3}$. These are lower than that of the rich Coma cluster and higher than that of the Virgo cluster.    Figure~\\ref{fdist1} shows the spatial distribution of the mergers visually identified among the sample of intermediate-mass ($M_{*} \\geq 10^{9} M_{\\sun}$) galaxies in the A901a, A901b, and A902 clusters.  We again consider here only the 20 distorted mergers listed in Table~\\ref{tmer}, for the reasons discussed in $\\S$Section~\\ref{sinte1}.  We find that all the 20 mergers lie outside the cluster cores.  We estimate that the number density ($n_{\\rm merge}$) of mergers is 0, 2.09, and 0.11 galaxies Mpc$^{-3}$, respectively, in the core, outer region, and outskirt of the cluster (Table~\\ref{t3rad}).  When estimating $n_{\\rm merge}$ in the three regions, we divided the number of mergers by the volume of the core region, and the volume of a spherical shell in the outer region and outskirt of the clusters.  For mergers and non-interacting galaxies, we compute the minimum distance to the nearest cluster center (A901a, A901b, and A902), and the local values of various environmental parameters, such as the local galaxy surface density $\\Sigma_{\\rm 10}$, the local DM mass surface density $\\kappa$, and the relative local ICM density.  Following Wolf \\etal (2005) and Gilmour \\etal (2007), we define and compute the local galaxy surface density ($\\Sigma_{\\rm 10}$) as the number of galaxies per (Mpc/h)$^{-2}$ that are in an annulus bracketed by two circles whose radii are equal to the average distances to the 9th and 10th nearest neighbor.  The DM mass surface density ($\\kappa$) is computed locally using the coordinates of each galaxy on the weak lensing map provided by Heymans \\etal (2008). The relative local ICM counts are measured at each galaxy position using the X-ray map from Gray \\etal (2010, in preparation) and Gilmour \\etal (2007).  Figure~\\ref{fnsigk} shows these local environmental parameters as a function of the minimum distance to the cluster center for mergers and non-interacting galaxies.  Since the mergers lie outside the cluster core (0.25 Mpc~$< R \\le$~2 Mpc), they are associated with low values of $\\kappa$ and intermediate values of $\\Sigma_{\\rm 10}$, and ICM density (Figure~\\ref{fnsigk}).   \\subsection{What accounts for the distribution of mergers?}\\label{sinte3}  The timescale for collisions and close encounters is given by: \\begin{equation} t_{\\rm coll} = \\frac{1}{n \\sigma_{\\rm gal} A}, \\end{equation}  where $n$ is the galaxy number density, $\\sigma_{\\rm gal}$ is the galaxy velocity dispersion, $A$ is the collisional cross section ($\\pi f(2r_{\\rm gal})^{2}$), $r_{gal}$ is a typical galaxy radius, and $f$ is the gravitational focusing factor (Binney \\& Tremaine, 1987). Since the average velocity dispersion in clusters is high, the value for $f$ will be low and we will assume it to be on the order of unity. Thus,  \\begin{equation}\\label{tcoll} \\begin{split} t_{\\rm coll}\\sim800\\! \\times \\left(\\!\\frac{n}{10^{3} \\rm \\, {\\rm  Mpc}^{-3}}\\!\\right)^{-1}\\! \\times \\left(\\frac{\\sigma}{10^{3}\\,{\\rm km\\,s}^{-1}}\\right)^{-1}\\! \\\\ \\times \\left(\\!\\frac{{r_{\\rm gal}}}{10 \\, {{\\rm kpc}}}\\!\\right )^{-2}\\! {\\rm Myr} \\end{split} \\end{equation}   We compute $n$ using the same method as in $\\S$Section~\\ref{sinte2} for the core, and use spherical shells to compute $n$ in the outer region and cluster outskirt for our sample of $R_{\\rm Vega} \\le$~24, intermediate mass ($M_{*} \\geq10^{9} M_{\\sun}$) galaxies. The local galaxy velocity dispersion profiles for A901a/b, A902, and the South West Group (SWG) from kinematic modeling using $\\sim$420 2dF redshifts (Gray \\etal 2010, in preparation) are shown in Figure~\\ref{fcvdp}, and the average velocity dispersions $\\sigma_{\\rm gal}$ are shown in Table~\\ref{tcompa}.  The central galaxy velocity dispersion within the cores ($R<0.25$~Mpc) of A901a,b and A902 typically range from 700 to 1000 km s$^{-1}$. Beyond the cluster core, in the outer region (0.25 Mpc~$< R \\le$~1.2 Mpc), the small number statistics leads to large error bars on the galaxy velocity dispersion, making it not viable to determine whether it remains high, drops, or rises. To estimate timescales, we take the average $\\sigma_{\\rm gal}$ for A901/902 to be $\\sim$800 km s$^{-1}$ in three regions of the cluster.  The timescales ($t_{\\rm coll}$) for close encounters are shown in Table~\\ref{t3rad} and range from 0.765 to 13.5 to 135 Gyrs from the core ($R\\le$~0.25 Mpc), to the outer region (0.25 Mpc~$< R \\le$~1.2 Mpc), to the outskirt region (1.2 Mpc~$< R \\le$~2.0 Mpc) of the cluster.   It may at first seem surprising that mergers do not preferentially lie in the cluster core, where the probability for close encounters is high due to the associated short value of the timescale $t_{\\rm coll}$ for collisions and close encounters. However, if $\\sigma_{\\rm gal}$ is large, a close encounter is unlikely to lead to a merger or a large amount of tidal damage (i.e. to strong morphological and kinematical distortions).  This is because a large galaxy velocity dispersion will likely cause the energy {\\it E} of the reduced particle in a binary encounter to be positive, causing the orbit to become unbound (Binney \\& Tremaine, 1987).  All merging systems lie at a projected clustocentric radius of 0.25 Mpc~$<R \\le$~2 Mpc, between the core and cluster outskirts, although in this region the timescale for collisions and close encounters is quite large.  We suggest two possible reasons why strong interactions and mergers might populate the outer region and cluster outskirt. The first possibility is that the velocity dispersion of galaxies falls between the core and outskirt, due to the {\\it intrinsic} velocity dispersion profile of the clusters.  Many clusters show a flat velocity dispersion profile from 1 Mpc outward, but several also show a declining profile (den Hartog \\& Katgert 1996, Carlberg 1997). In the latter cases, the velocity dispersions can fall from 1500 to 100 km s$^{-1}$ from the cluster core to cluster outskirt.  Since the local velocity dispersion profiles in the regions of interest in A901/902 (Figure~\\ref{fcvdp}) are based on a small number of galaxies and therefore large errors in $\\sigma_{\\rm gal}$, we cannot determine if the velocity dispersion profile falls or remains flat.  An additional difficulty is that the member galaxies of the neighboring clusters influence $\\sigma_{\\rm gal}$ at $\\sim$0.5 Mpc for A901a/b and $\\sim$1 Mpc for A902.  The second possibility is that the mergers are part of groups or field galaxies that are being accreted along filaments by the A901/902 clusters.  This would be in line with the idea that clusters continuously grow by mergers and accretion of groups (e.g., Zabludoff \\& Franx 1993; Abraham \\etal 1996a; Balogh, Navarro \\& Morris 2000). Groups have lower galaxy velocity dispersion (typically below 300 km s$^{-1}$) than clusters, as shown in Table~\\ref{t3rad}.  Furthermore simulations show that merger rates are highest as field and group galaxies accrete into cluster along cosmological filaments (see below; van Kampen \\& Katgert 1997; Figure~\\ref{felco1}).   To further quantify the possibility of groups accreting into the A901/A902 clusters, we compare the data to model predictions from simulations of the STAGES A901/A902 supercluster (van Kampen \\etal in preparation).  The simulation aims to reproduce the environment as observed for A901/A902 by using constrained initial conditions that produce clusters with overall properties similar to those of A901, A902, and the neighboring A907 and A868.  The simulation box is large enough to contain all these, and hence reproduce the overall large-scale density and velocity fields similar to what is observed. A phenomenological galaxy formation model is then used to simulate the galaxy population in the same box, allowing us to select mock galaxies above a certain luminosity or stellar mass cut, as done in the observed dataset.  A stellar mass cut of $M_{*}\\geq$ 10$^{9}$ $M_{\\sun}$ is applied in the simulations, corresponding to the mass cut in the observational sample. The solid lines in Figure~\\ref{felco1} show the predicted number density ($n$) and fraction ($f$) of galaxy mergers with mass ratio $M_{1}/M_{2} >$~1/10, as a function of environment, as characterized by the local overdensity ($\\delta^{\\rm G}$).  The latter is calculated by smoothing the density of DM halos with a Gaussian of width 0.4 Mpc to take out the effect of individual galaxies.  Typical values of $\\delta^{\\rm G}$ are $\\sim$10-100 for group overdensities, $\\sim$200 at the cluster virial radius, and $\\gtrsim$~1000 in the core of rich clusters.  Typically, in the simulations, as field and group galaxies fall into a cluster along filaments, the coherent bulk flow enhances the galaxy density and causes galaxies to have small relative velocities, thus leading to a high probability for mergers at typical group overdensities.  Closer to the cluster core, galaxies show large random motions rather than bulk flow motion, leading to a large galaxy velocity dispersion, and a sharp drop in the probability of mergers.   In order to compare the data to simulation results, we use the number ($N_{\\rm merge}$) and fraction ($f_{\\rm merge}$) of mergers of mass ratio $\\ge 1/10$, which we computed in $\\S$Section~\\ref{sinte1}.  The solid curve in Figure~\\ref{felco1} shows the merger fraction and number density predicted by the simulations as a function of overdensity. The three dashed lines in Figure~\\ref{felco1} show the estimated range in the observational number density ($n_{\\rm merge}$) and fraction ($f_{\\rm merge}$) of mergers in the three different regions (the core, the outer region, and the outskirt) of the A901/902 clusters, as defined in $\\S$Section~\\ref{sinte2}.  The points at which the dashed lines cross or approach the solid curve tell us the typical overdensities at which we expect to find such merger fraction or merger number densities in the simulations.  It can be seen that the low merger density seen in the core region of the cluster correspond to those expected at typical cluster core overdensities.  On the other hand, the larger merger fraction we observe between the cluster core and outer region (0.25 Mpc~$< R \\le$~2 Mpc) is close to those seen in typical {\\it group overdensities}.  Our results are in agreement with the above scenario, whereby the accretion of group and/or field galaxies along filaments, where they have low relative velocities and enhanced overdensities, leads to enhanced merger rates.  We note that similar conclusions are reached if we perform the comparisons between data and simulations using the more conservative mass cut of $M_{*} \\geq 3 \\times 10^{9} M_{\\sun}$, discussed in $\\S$Section~\\ref{sinte1}.  However, there are a couple of caveats when comparing our data with simulations. One caveat is that we directly compare the projected values of the number density or fraction of mergers from two dimensional observations, to the predicted number density of mergers from three dimensional simulations.  Projection effects can introduce uncertainties in our observational estimate. The magnitude of the uncertainties due to projection effects depends on the detailed environment and will be investigated in the full STAGES supercluster simulation dataset (van Kampen \\etal 2009, in preparation).   A further indirect evidence for group accretion stems from comparing semi-analytic galaxy catalogs to STAGES observations. Rhodes \\etal (in preparation) finds an overabundance in galaxies in A902 compared to its mass. This could be explained by  two or more galaxy groups in projection, consistent with the idea of group accretion.   \\subsection{Comparison with groups and clusters at different epochs}\\label{scomp1}  We first recapitulate our results on the visually-based merger fraction ($f_{\\rm merge}$) in the A901/902 clusters ($\\S$Section~\\ref{sinte1}).  Among intermediate-mass ($M_{*} \\geq10^{9} M_{\\sun}$) systems, we find that the fraction $f_{\\rm merge}$ of systems which show evidence of a recent or ongoing merger of mass ratio $ >$~1/10 is 0.023$\\pm$0.007 (Table~\\ref{tcut}).  Most of these mergers have $M_{\\rm V} \\sim -19$ to $-22$ and $L \\le L^{*}$ (see Figure~\\ref{flumf}).  We also estimated that the fraction of likely major mergers, likely minor mergers, and ambiguous cases to be 0.01$\\pm$0.004\\% (9/886), 0.006$\\pm$0.003\\% (5/886), and 0.007$\\pm$0.003\\% (6/886), respectively.  Next, we compare our merger fraction in the A901/902 clusters with the published merger fraction in other clusters and group galaxies out to $z\\sim$~0.8, over the last 7 Gyr (Figure~\\ref{fcompa}). These comparisons are difficult to make as different studies apply different luminosity or mass cuts.  Furthermore, some studies consider only major mergers, while others consider all interacting galaxies, which are likely candidates for both major and minor interactions.  The variation in galaxy populations sampled at different epochs must also be kept in mind when comparing results at lower redshift with those out to $z\\sim$~0.8.  Low redshift samples typically sample a small volume and therefore tend to host only a small number of the brightest and most massive systems.  Conversely, higher redshift magnitude-limited samples will suffer from Malmquist bias and tend to under-represent the faint ($L < L^{*}$) galaxies.  There have been several studies of galaxy mergers and interactions in intermediate redshift clusters (Lavery \\& Henry 1988; Lavery, Pierce, \\& Mclure 1992; Dressler \\etal 1994; Oemler, Dressler, \\& Butcher 1997; Couch \\etal 1998), but only few to date with high resolution $HST$ imaging.  Couch \\etal (1998) used WFPC/WFPC2 observations and spectroscopy of two clusters at $z\\sim$0.3 and $\\sim0.4$, focusing on galaxies with $R\\sim$22.5 or $L < L^{*}$.  In their study, they classified $\\sim$7/200 or $\\sim3.5\\%$ galaxies to be merging based on separations and visible distortions. The merger fractions of these two intermediate redshift clusters are plotted in Figure~\\ref{fcompa}.  It is also interesting to note that the merging systems in these two intermediate redshift clusters tend to be blue and preferentially located in the outskirt of the clusters.  In fact, in the Couch \\etal (1998) sample, $\\sim60\\%$ of the merging galaxies among $L <L^{*}$ systems are blue.  The mergers in these two intermediate redshift clusters are similar in many ways to the mergers in the A901/902 clusters: the latter mergers lie in between the cluster core and outskirt (0.25 Mpc~$< R \\le$~2 Mpc; $\\S$Section~\\ref{sinte2}) of the A901/902 clusters, and for the absolute magnitude range of $M_{\\rm V} \\sim$-19 to -22, 16/20 or $80 \\pm 18\\%$ of these mergers lie on the blue cloud ($\\S$Section~\\ref{scolor1}).   In a study by Dressler \\etal (1994), based on WFPC images of a $z\\sim0.4$ cluster, the fraction of bright ($M_{V} < -18.5$) galaxies identified as mergers was 10/135.  Another WFPC study of four clusters at $z=0.37-0.41$ by Oemler, Dressler, \\& Butcher (1997) placed a lower limit of 5\\% on the fraction of merging galaxies with $M_{V} < -19$ or $L \\leq L^{*}$.   A high redshift ($z=0.83$) cluster was studied by van Dokkum \\etal (1999) using $HST$ WFPC2 images and spectroscopic redshifts for 80 confirmed members.  They found a high fraction (17\\%) of luminous ($M_{B}\\sim -22$; $L \\sim 2L^{*}$) systems, which exhibit signatures of ongoing mergers.  These mergers were found to be red, bulge-dominated, and located in the outskirts of the cluster.  A later study by Tran \\etal (2005b) used spectroscopic followup to confirm that 15.7\\%$\\pm3.6\\%$ of these galaxies were bound pairs with separations $\\ltsim$30h$^{-1}$ kpc and relative velocities $\\ltsim$300 km s$^{-1}$. These merging pairs are located outside the cluster centers, similar to the mergers in A901/902. However, the galaxy sample at high redshift is dominated by more luminous galaxies ($M_{B}\\sim$ -22; $L \\sim2L^{*}$) than the A901/902 sample where most galaxies have $M_{V}\\sim$ $-$19.5 to $-$22 (Figure~\\ref{flumf}).  The number statistics for bright ($\\sim 2L^{*}$) mergers in A901/902 are not robust enough to allow a direct comparison with the merger fraction cited in van Dokkum \\etal (1999).  Finally, we look at the merger fraction in {\\it groups}.  A study by McIntosh \\etal (2008) finds that $\\sim$1.5\\% of massive ($M \\ge 5 \\times 10^{10} M_{\\sun}$) galaxies in SDSS groups (with $M_{\\rm halo} > 2.5 \\times 10^{13}$ $M_{\\sun}$) are major mergers at $z$ of 0.01--0.12.  Most of the mergers involve two red sequence galaxies and they are located between 0.2 and 0.5 Mpc from the group center.  A study by McGee \\etal (2008) also finds an enhancement in asymmetric bulge-dominated galaxies in groups, consistent with a higher probability for merging in the group environment.  Additional evidence for mergers in groups is shown in a study of a supergroup at $z\\sim0.37$ by Tran \\etal (2008), who report dry dissipationless mergers and signatures thereof in three of four brightest group galaxies.   A merger fraction of 6\\% was found by Zepf (1993) in Hickson compact groups at $z< 0.05$, among systems with luminosities $L \\leq L^{*}$. The merger fraction is significantly higher than that in SDSS groups. The increased merger fraction is expected since Hickson compact groups are different from loose groups: they have high number densities comparable to those in cluster cores, but low galaxy velocity dispersions. These two conditions provide an environment most favorable to strong tidal interactions and mergers, as argued in $\\S$Section~\\ref{sinte3}.   It is not straightforward to draw conclusions about the evolution of the merger fraction in cluster galaxies over $z\\sim$~0.05--1.0 from the above studies because they sample different types of systems and different luminosity ranges. If we conservatively restrict ourselves to only studies of $L \\leq L^{*}$ cluster galaxies, then we have 4 data points over $z\\sim$~0.17--0.4, shown as solid filled circles in Figure~\\ref{fcompa}.  The mean value of the data points allows for evolution by a factor of $\\sim$3.2, in the merger fraction of $L \\leq L^{*}$ cluster galaxies over $z\\sim$~0.17--0.4.  However, if one uses the full range of merger fractions allowed by the error bars on the data points, then we can admit a wider spectrum of scenarios, ranging from no evolution to evolution by a factor of $\\sim$5 over $z\\sim$~0.17--0.4.  Having additional deep, large-volume, high resolution studies, which are based on larger samples with smaller error bars would help to separate between these scenarios, and thereby test hierarchical models of galaxy evolution.  As mentioned in $\\S$Section~\\ref{sinte1}, hierarchical models (e.g., Gottloeber et al 2001; Khochfar \\& Burkert 2001) predict that the merger fraction in dense clusters falls more steeply at $z<1$ than the field merger fraction, such that at $z<0.3$, the predicted merger fraction is quite low (typically below 5\\%) among intermediate mass cluster galaxies.  The low merger fractions among intermediate mass ($10^{9}$ to a few $\\times 10^{10} M_{\\sun}$) or intermediate luminosities ($L < L^{*}$) galaxies in the A901/902 clusters and other low redshift clusters are consistent with these predictions, but we cannot yet test the predicted rate of evolution of the merger fraction in clusters with redshfit.     \\subsection{Galaxies on the blue cloud and red sample}\\label{scolor1}   We explore the properties of galaxies on the blue cloud and in the red sample among the final sample of 886 systems (20 mergers and 866 non-interacting galaxies; $\\S$Section~\\ref{smvc}). The red sample was defined in Wolf \\etal (2009) and contains both passively evolving spheroidal galaxies on the red sequence, as well as dusty red galaxies that lie between the red sequence and the blue cloud.  Figure~\\ref{fcmd1} shows the rest-frame $U\\!-\\!V$ color plotted against stellar masses for galaxies of different visual classes: Mergers, Non-interacting Irr-1, and Non-interacting Symmetric ($\\S$Section~\\ref{smvc}). The visual classes of galaxies on the blue cloud and in the red sample are shown in Table~\\ref{tvccol}.  As described in $\\S$Section~\\ref{sinte1}, we only consider the 20 distorted mergers listed in Table~\\ref{tmer}, and avoid the potential projection pairs without signs of morphological distortions.  The 20 mergers are divided into 13 mergers of type 1, 3 mergers of type 2a, and 4 mergers of type 2b. For mergers of type 2b, which are resolved into two galaxies with separate COMBO-17 colors, we plot the average $U-V$ color of the galaxies in the pair.  Out of the 886 visually classified systems in our sample, we find that 310/886 or $35\\%\\pm7\\%$ lie on the blue cloud.  Out of the 20 visually classified mergers with distortions, 16/20 or $80\\%\\pm 18\\%$ lie on the blue cloud (Table~\\ref{tvccol}).  Conversely, out of 866 non-interacting galaxies, 294/866 or $34\\%\\pm 7\\%$ are on the blue cloud.  Thus, the fraction of mergers, which lie on the blue cloud ($f_{\\rm merg-BC}$) is over two times larger than the fraction of non-interacting galaxies ($f_{\\rm non-int-BC}$), which lie on the blue cloud. This implies that mergers and interacting galaxies are preferentially blue, compared to non-interacting galaxies .   The observed higher fraction of blue galaxies among mergers compared to non-interacting galaxies may be caused by several factors. It may be due to enhanced levels of unobscured SF in mergers (see $\\S$Section~\\ref{ssfr1}), translating to bluer colors on average. It may also in part be the result of the mergers being part of accreted field and group galaxies ($\\S$Section~\\ref{sinte3}), which are bluer than the average cluster galaxy.  It is also relevant to ask whether we are overestimating the fraction of blue galaxies among interacting systems due to the visibility timescale $t_{\\rm vis}$ of morphological distortions induced by interactions being longer in bluer galaxies than redder ones.  While this is possible, it is non-trivial to correct for this effect because no direct unique relation exists between $t_{\\rm vis}$ and color.  As discussed in $\\S~\\ref{sinte1}$, $t_{\\rm vis}$ depends on the mass ratio of an interaction as well as the gas content.  A higher gas content may result in a longer $t_{\\rm vis}$ for certain gas distributions, but it can lead to either redder colors (e.g., enhanced level of dusty SF) or bluer colors (enhanced level of unobscured SF), compared to interacting systems.    \\subsection{SF properties of interacting galaxies}\\label{ssfr1} This work uses SFRs based on UV data from COMBO-17 (Wolf \\etal 2004) and $Spitzer$ 24$\\micron$ imaging (Bell \\etal 2007). The unobscured SFR$_{\\rm UV}$ is derived using the 2800\\ \\AA \\ rest frame luminosity ($L_{\\rm UV}$ = 1.5$\\nu l_{\\nu,2800}$) as described in Bell \\etal (2005, 2007).  The UV spectrum is dominated by continuum emission from massive stars and provides a good estimate of the integrated SFR from the younger stellar population in the wavelength range of 1216-3000\\ \\AA. \\ The SFR$_{\\rm IR}$ is derived using the 24$\\micron$ flux to construct the integrated IR luminosity ($L_{\\rm IR}$) over 8-1000$\\micron$ following the methods of Papovich \\& Bell (2002).  The total SFR is derived using identical assumptions of Kennicutt (1998) from PEGASE assuming a 100 Myr old stellar population with constant SFR and a Chabrier (2003) IMF: \\begin{equation}\\label{sfrtot} {\\rm SFR}_{\\rm UV + IR} = 9.8 \\times 10^{-11}(L_{\\rm IR} + 2.2L_{\\rm UV}). \\end{equation} The factor of 2.2 on the UV luminosity accounts for light being emitted longward of 3000\\ \\AA \\ and shortward of 1216\\ \\AA \\ by young stars.  The total SFR accounts for both the dust-reprocessed (IR) and unobscured (UV) SF.   We work with our final sample of 886 classifiable bright massive ($M_{*} \\geq 10^{9} M_{\\sun}$) systems, and only consider here the 20 distorted mergers in Table~\\ref{tmer}.  All of these systems have UV-based SFR from COMBO-17 observations. Of this sample, $\\sim$11\\% (94/886) were not observed with $Spitzer$, $\\sim$23\\% (206/886) were observed and detected at 24$\\micron$, while the rest had no detection at the $\\sim$4$\\sigma$ depth of 83 $\\mu$Jy.   The UV-based SFR (SFR$_{\\rm UV}$) versus stellar mass is plotted in Figure~\\ref{fsfruv} for all 886 systems.  The SFR$_{\\rm UV}$ ranges from $\\sim$0.01 to 14 $M_{\\sun}$ yr$^{-1}$.  The UV-based SFR represents a lower limit to the total SFR for galaxies on the blue cloud and most star-forming galaxies on the red sample. However, for some old red galaxies, the SFR$_{\\rm UV}$ may overestimate the true SFR as the UV light from such systems may not trace massive OB stars, but rather low to intermediate mass stars.  Figure~\\ref{fsfrtot} shows the UV+IR-based SFR (SFR$_{\\rm UV + IR}$), which ranges from $\\sim$0.2 to 9 $M_{\\sun}$ yr$^{-1}$.  For the 206 galaxies that were observed and detected at 24$\\micron$, the implied UV+IR-based SFR is plotted as stars in the lower panel of Figure~\\ref{fsfrtot}.  For the 586 galaxies that are observed but undetected with $Spitzer$, we use the 24$\\micron$ detection limit as an upper limit on the 24$\\micron$ flux.  The corresponding upper limit on the UV+IR-based SFR is plotted as inverted triangles in the lower panel of Figure~\\ref{fsfrtot}, and included in the calculation of the average UV+IR-based SFR, plotted in the middle panel.   In a cluster environment, the competition between processes that enhance the SFR and those that depress the SFR, ultimately determine the average SFR of cluster galaxies.  The first class of processes are strong close interactions: tidal and mergers (e.g., Toomre \\& Toomre 1972), and harassment (e.g., Moore \\etal 1996), which refers to the cumulative effect of weak interactions.  The second class of processes include ram pressure stripping of cold gas out of the galaxy by the hot ICM (e.g., Gunn \\& Gott 1972), and strangulation (e.g., Larson \\etal 1980; Balogh, Navarro \\& Morris 2000).  For the few mergers (orange line) present, the average SFR$_{\\rm UV}$ is typically enhanced by an average modest factor of $\\sim$2 compared to the both Non-interacting Symmetric and Irr-1 galaxies (purple, green and black lines).  Similarly, the UV+IR-based SFR (SFR$_{\\rm UV + IR}$; and Figure~\\ref{fsfrtot}) of merging galaxies is typically enhanced by only an average factor of $\\sim$1.5 compared to the Non-interacting Symmetric galaxies (purple line) and to all Non-interacting galaxies (i.e Symmetric + Irr-1; black line).  We note that a similar modest enhancement in the average SFR, by a factor of 1.5--2 is also found in mergers in the field over $z\\sim$~0.24--0.80 by Jogee \\etal (2008, 2009). This modest enhancement is consistent with the theoretical predictions of di Matteo \\etal (2007; see their Figure 10), based on a recent statistical study of several hundred simulated galaxy collisions.  Modest SFR enhancements are also seen in galaxy pair studies in the field (Barton \\etal 2000,2003; Lin \\etal 2004; Ellison \\etal 2008) and in mixed environments (Robaina \\etal 2009; Alonso \\etal 2004).  While mergers in the A901/902 clusters enhance the SFR of individual galaxies, it is clear that they do not contribute much to the total SFR of the cluster. We compute the SFR density of mergers to non-interacting galaxies in the same volume of A901/902 by taking the ratio of the total SFR in each class.  If we include only the 20 distorted mergers in Table~\\ref{tmer}, we find that the contribution of mergers to the SFR density of the clusters to be 10\\%. Alternatively, if we include all of the 36 visually classified mergers in Table~\\ref{tmer}, before accounting for false projection pairs, we find the contribution of mergers to the SFR density to be 15\\%. Thus, we find that mergers contribute only a small fraction (between 10\\% and 15\\%) of the total SFR density of the A901/902 clusters compared to non-interacting galaxies. The small contribution of mergers to the total cluster SFR density is likely due to the low number of mergers in A901/902, and the fact that these mergers only cause a modest SFR enhancement."
    },
    "0910/0910.0883_arXiv.txt": {
        "abstract": "We describe two-dimensional gasdynamical computations of the X-ray emitting gas in the rotating elliptical galaxy NGC 4649 that indicate an inflow of $\\sim1$ $M_{\\odot}$ yr$^{-1}$ at every radius.  Such a large instantaneous inflow cannot have persisted over a Hubble time. The central constant-entropy temperature peak recently observed in the innermost 150 parsecs is explained by compressive heating as gas flows toward the central massive black hole.  Since the cooling time of this gas is only a few million years, NGC 4649 provides the most acutely concentrated known example of the cooling flow problem in which the time-integrated apparent mass that has flowed into the galactic core exceeds the total mass observed there.  This paradox can be resolved by intermittent outflows of energy or mass driven by accretion energy released near the black hole.  Inflowing gas is also required at intermediate kpc radii to explain the ellipticity of X-ray isophotes due to spin-up by mass ejected by stars that rotate with the galaxy and to explain local density and temperature profiles.  We provide evidence that many luminous elliptical galaxies undergo similar inflow spin-up.  A small turbulent viscosity is required in NGC 4649 to avoid forming large X-ray luminous disks that are not observed, but the turbulent pressure is small and does not interfere with mass determinations that assume hydrostatic equilibrium. ",
        "introduction": "X-ray isophotes of massive, rotating elliptical galaxies convey valuable information about the rotation of the hot, virialized interstellar gas, the interaction of this gas with gas recently ejected from evolving stars, and the sense of radial motion -- in or out -- of the hot gas. These issues are relevant to understanding the physical nature of the hot gas in galaxy groups and clusters in which gas loses energy by radiating X-rays but does not appear to cool to low temperatures as expected. In traditional cooling flows the radiating gas approximately maintains the local virial temperature as it flows slowly toward the center of the confining gravitational potential, then cools catastrophically near the center. The so-called cooling flow problem arises because there is insufficient spectroscopic evidence for cooling gas, previously cooled gas or recently formed stars. This paradox is apparent in massive galaxy clusters, galaxy groups and individual elliptical galaxies, i.e. the cooling flow problem is independent of scale. For this reason observations of X-ray luminous, relatively nearby elliptical galaxies can provide important clues to unravel the mysteries of cooling flows that do not seem to cool.  The recent deep 81ks Chandra observation of the nearby X-ray luminous elliptical galaxy NGC 4649 by Humphrey et al. (2008) led to the detection of a massive central black hole ($M_{bh} = 3 \\times 10^{9}$ $M_{\\odot}$). This supermassive black hole was found for the first time directly from X-ray data by assuming that the hot gas is very nearly in hydrostatic equilibrium. The increased quality of the X-ray observations of NGC 4649 has motivated us to develop a dynamical model for the hot gas flow in this rotating galaxy.  At any time the radial flow of hot gas must either be inward, outward, a combination of the two, or stationary. Many people think that cooling inflows cannot occur in elliptical galaxies or elsewhere because of the absence of X-ray spectral emission at low and intermediate temperatures (e.g. Xu et al. 2002). We showed for galactic-scale cooling flows that dust expelled from stars can cool the gas so fast that very little low-temperature thermal X-ray emission occurs (Mathews \\& Brighenti 2003a). Nevertheless, cooling inflows cannot dominate over time since the central mass accumulated would greatly exceed limits set by the velocity dispersion of central stars. Detailed gasdynamical models with {\\it ad hoc} heating sources have shown that the transition from cooling inflows to strong low-density wind outflows is very abrupt (Mathews \\& Brighenti 2003b). In the absence of extreme fine tuning, most of the X-ray luminous hot gas bound to elliptical galaxies cannot be flowing globally outward, although inhomogeneous outflowing regions of lower gas density (and X-ray emissivity) in jets or buoyant regions are not ruled out. If the gas is not flowing at all, continued enrichment by Type Ia supernovae for a few Gyrs will raise the iron abundance to several times that of the underlying stars, which is not generally observed (e.g. Humphrey \\& Buote 2006). One way to avoid this enrichment catastrophe is to recognize that some or most of the SNIa iron may cool radiatively in a way that cannot be observed, as discussed by Brighenti \\& Mathews (2005). However, the modest SNIa enrichment observed in the hot gas can be understood if much of the hot galactic gas flows radially inward, and the abundance can be even further reduced if more extended circumgalactic group-scale gas with relatively low metallicity flows inward past the sources of stellar enrichment. Circumgalactic gas is likely since Humphrey et al. (2006) derive a group-level virial mass for NGC 4649, $3.5 \\times 10^{13}$ $M_{\\odot}$.  Several additional observations of NGC 4649 can be understood most easily in terms of gaseous inflow. First, the central gas temperature peak discovered by Humphrey et al. (2008) within a few 100 parsecs was predicted by cooling {\\it inflow} models as gas approaches a massive central black hole (Brighenti \\& Mathews 1999). While only a central gas pressure peak is sufficient to determine the black hole mass, a peak in the gas temperature is a signature of compressional heating as hot gas flows inward toward a strongly concentrated mass or small rotating disk. This compressive heating can occur even when the gas is losing energy by standard radiative losses. Additional support for subsonic compressional heating during inflow follows from the observation that the entropy of hot gas closest to the black hole is lower than in any other region in NGC 4649 so it has not been recently heated by the central AGN.  Another important signature of subsonically inflowing gas seen in NGC 4649 and other similar elliptical galaxies is a central flattening of X-ray isophotes on kpc scales. Elliptical X-ray isophotes on galactic scales have different implications than on larger scales which have been used to infer flattening of the dark halo potential (Buote et al. 2002). Isophotal flattening in excess of that required by the underlying stellar potential can be interpreted as a spinning up of hot inflowing gas expected when angular momentum is exchanged with gas recently ejected from old, mass-losing red giant stars that participate in the galactic rotation. We review below the recent compilation of Chandra isophotes of elliptical galaxies by Diehl \\& Statler (2007) and show that they generally resemble NGC 4649 in central flattening and implied inflow spin-up. NGC 4649 is an ideal candidate galaxy to study the influence of rotation on the X-ray isophotes because of its remarkably high rotation rate for a massive, cored E galaxy, its proximity, and our good fortune to have received at least one moderately deep Chandra observation.  In the following discussion we review the implications of recent Chandra images of NGC 4649 and present gasdynamical models of increasing complexity. Unlike our previous studies of rotating cooling flows (Brighenti \\& Mathews 1996, 2000), we include the important contribution of inflowing and non-rotating circumgalactic gas from a group-scale halo. In addition, since the central X-ray isophotal flattening observed in NGC 4649 and other massive ellipticals is less than that expected from the conservation of angular momentum, we invoke an {\\it ad hoc} hot gas turbulence that can transport angular momentum away from the galactic spin axis. By combining turbulence and circumgalactic inflow it is easy to explain the isophotal flattening observed in NGC 4649. Moreover, the turbulent viscosity that is consistent with the X-ray isophotes of NGC 4649 generates a turbulent pressure that is much less than the gas pressure. Therefore, the subsonic turbulent activity required to reduce the angular momentum does not significantly degrade the total mass determinations of Humphrey et al. (2008) based on assuming hydrostatic equilibrium. We assume that the mild turbulence we require is produced by energy released near the central black hole or active galactic nucleus (AGN).  Our treatment here is done under the assumption that most of the {\\it observed} X-ray emission in NGC 4649 can be explained with inflowing gas. We do not propose that this is the complete solution to the gas flow problem in this or other similar galaxies, but only demonstrate that the X-ray observations are consistent with an {\\it apparent} global inflow at the present time. To avoid unobserved catastrophic radiative cooling and a huge mass concentration in the galactic core, we must also assume that there is a (possibly buoyant) return mass outflow that is too intermittent, too hot, or too cunningly disguised to contribute to Chandra observations (Mathews \\& Brighenti 2004; Brighenti \\& Mathews 2006). ",
        "conclusions": "We have shown that the X-ray emitting gas associated with the bright elliptical galaxy NGC 4649 has density and temperature profiles that can be understood as a rotating, radiatively cooling inflow with a mild turbulent viscosity. Evidence of inflow is present at every radius in NGC 4649. On the largest scales, a radiatively cooling inflow of circumgalactic gas is required to dilute the hot gas iron abundance acquired from supernovae down to observed values, assuming that the iron ejected by Type Ia supernovae does not cool by radiative losses before it merges with the ambient hot gas. The relatively high gas density and temperature observed in NGC 4649 beyond $\\sim$10 kpc can also be explained with inflowing circumgalactic gas. On intermediate galactic scales, observations of the X-ray ellipticity in NGC 4649 and other bright elliptical galaxies show that the hot gas is spun up by mass ejected from evolving stars that rotate collectively with the galaxy. Inflow spin-up is seen in the increasing X-ray isophotal ellipticity with decreasing galactic radius until it exceeds the ellipticity of the local stars near the galactic center. However, for NGC 4649, which rotates unusually fast for a massive boxy elliptical, a small turbulent viscosity is required to avoid forming multi-kpc X-ray disks that are not observed. The shallow negative temperature gradient inside $\\sim1$ kpc in NGC 4649 is X-ray evidence of an inflow that compresses toward a sub-kpc disk. The observed X-ray isophotes can be matched with a turbulent viscosity for which the corresponding turbulent pressure is much less than the gas pressure, so the integrated mass profiles found by assuming hydrostatic equilibrium (Humphrey et al. 2008) are unaffected. This alleviates the concern expressed by Diehl \\& Statler (2007) that the hot gas in elliptical galaxies is very far from hydrostatic equilibrium. Galactic mass determinations based on X-ray thermal emission may fail for some elliptical galaxies where non-thermal pressure dominates the galactic core, such as NGC 4636 (Brighenti \\& Mathews 1997; Baldi et al. 2009), but this problem does not seem to occur in NGC 4649. On the smallest scales observed with Chandra, within about 100 parsecs, the central temperature peak in NGC 4649 discovered by Humphrey et al. (2008) is a natural consequence of subsonic cooling inflow. In spite of ongoing radiative losses, the central gas temperature increases due to (nearly adiabatic) compression as gas approaches the central black hole or the small disk around it. The kpc-scale negative temperature gradient in NGC 4649 formed as gas compresses toward a small disk is not caused by recent AGN heating since the gas entropy in this region increases with galactic radius. This type of central inflow is also consistent with nuclear disks of cooled dusty gas having radii of a few hundred parsecs commonly observed in the cores of luminous elliptical galaxies.  NGC 4649 represents the most concentrated known example of the cooling flow problem. The observed density and temperature profiles can be explained at every radius with a radiatively cooling inflow. At the smallest observable radius in NGC 4649, about 150 pc, the times for cooling and flow to the core are only a few million years. For this reason it will be very difficult to devise heating scenarios that maintain the gas at rest, perfectly mimicking a cooling inflow. No ideally heated model in which the hot gas in NGC 4649 remains stationary can account for commonly observed X-ray ellipticity profiles $\\epsilon_X(r)$ that require inflow spin-up.  The current mass inflow rate in our successful calculations is $\\sim1$ $M_{\\odot}$ yr$^{-1}$, but the stellar mass loss rate was certainly several times larger in the past. Consequently, if this flow continued over a Hubble time, the total mass of cooled gas accumulated in the galactic core of NGC 4649 would exceed the mass of the black hole and stars within $\\sim1$ kpc by factors of 3 - 10. Therefore, the mass inflow indicated by the central temperature peak must be removed from within $\\sim150$ pc at a (time-averaged) rate comparable to $\\sim1$ $M_{\\odot}$ yr$^{-1}$ and transported out to a large radius without interfering with the X-ray appearance of NGC 4649 which resembles a traditional cooling flow. This mass outflow may occur in the bipolar jets observed in NGC 4649 and many other bright ellipticals. Alternatively, nuclear hot gas may become buoyant by intermittent local heating or by cosmic rays and flow out subsonically upstream in the cooling flow (e.g. Mathews \\& Brighenti 2008). If the buoyant gas density is only slightly less than that of the ambient gas, it may not significantly disturb the radial gas density and temperature profiles set by the inflowing gas. Alternatively, if the density of buoyant (or jet) gas were considerably lower, its X-ray emission would be less easily observed against the brighter emission from denser inflowing gas.     \\vskip.1in"
    },
    "0910/0910.5142_arXiv.txt": {
        "abstract": "Cosmology is operating now on a well established and tightly constraining empirical basis. The relativistic $\\Lambda$CDM hot big bang theory is consistent with all the present tests; it has become the benchmark. But the many open issues in this subject make it reasonable to expect that a more accurate cosmology will have more interesting physics in the invisible sector of the universe, and maybe also in the visible part. ",
        "introduction": "\\begin{figure}[b] \\includegraphics[angle=90,height=.33   \\textheight]{cmb_spectrum.pdf} \\caption{Energy spectrum of the cosmic microwave background, from Alan Kogut.} \\end{figure}  Lest the exciting debates over the great open issues of cosmology cause us to forget, I offer in Figure~1 and Table~1 reminders that we have a firm and substantial empirical basis for our subject. The figure shows the spectrum of the cosmic microwave radiation that nearly uniformly fills space: energy density as a function of frequency. The measurements are diverse and well cross-checked, from the ground, balloons, rocket and satellite missions, and indirectly from spin temperatures derived from interstellar molecular absorption lines. Within the measurement uncertainties, which over much of the range of frequencies in this figure are smaller than the symbols, this wealth of data fits the Planck thermal spectrum at temperature $T=2.725$~K plotted as the solid curve.  We know the universe as it is now is close to transparent at the frequencies plotted here, because active galaxies at redshifts exceeding unity are detected at these frequencies. Relaxation to this thermal spectrum thus had to have happened earlier in the expansion of the universe, when the density of absorbing matter was great enough make the universe opaque across the Hubble length then. That is, the figure presents us with almost tangible evidence that the universe really did evolve from a very different state. We should take a moment to admire this wonderfully simple and deeply significant phenomenon.  The figure also shows that expansion and cooling to the present transparent state of the universe had to have preserved the thermal radiation spectrum formed when the universe was opaque. Sufficient conditions include standard local physics and a metric description of spacetime. It does not require general relativity theory. It does require that, despite the distinctly inhomogeneous distribution of matter in galaxies and concentrations of galaxies, spacetime is close enough to homogeneous and isotropic not to have produced a sigificant spread of redshifts along different lines of sight (as could have happened in a homogeneous anisotropic universe, for example). Space has to be close to what Einstein envisioned in 1917.  The $\\Lambda$CDM cosmology adopts Einstein's near homogeneous space, his general relativity theory, and his cosmological constant, and it adds the general expansion first discussed by Friedmann and Lema\\^\\i tre, Tolman's thermal radiation, Gamow's thermonuclear element formation, and the more recent notions of nonbaryonic dark matter and adiabatic near scale-invariant Gaussian initial departures from homogeneity. The application of general relativity is a long extrapolation from its precision tests, but it was natural to try this theory first. Other of these assumptions were not so obvious first guesses; they were adopted because they were seen to offer promising fits to the improving tests. We would be poor theorists indeed if $\\Lambda$CDM did not at least approximate what is now measured. But this cosmology has gone beyond that; it is predictive.    \\begin{table}[t] \\caption{Cosmological Tests} \\centering \\begin{tabular}{lllll} \\hline \\multicolumn{2}{l}{1. fossil CMB radiation} & \\multicolumn{2}{l}{6. cosmic  mass density} \\\\ & (a) energy spectrum &&  (a) galaxy peculiar velocities \\\\ & (b) temperature acoustic oscillation &&  (b) gravitational lensing \\\\ & (c) temperature-polarization cross spectrum & \\multicolumn{2}{l}{7. large-scale structure} \\\\ & (d) integral Sach-Wolfe effect && (a) baryon acoustic oscillation \\\\ \\multicolumn{2}{l} {2. light element abundances} && (b) galaxy N-point functions \\\\ & (a) deuterium in Ly$\\alpha$ absorbers at redshift $\\sim 3$ & \\multicolumn{2}{l}{8. clusters of galaxies}\\\\ & (b) helium in dwarf low-metallicity galaxies && (a) mass function as a function of mass \\& redshift \\\\ & (c) observed baryon budget && (b) baryon mass fraction \\\\ & (d) count of neutrino families & \\multicolumn{2}{l}{9. small-scale structure}\\\\ \\multicolumn{2}{l}{3. redshift-magnitude relation} && (a) galaxy structure and evolution \\\\ \\multicolumn{2}{l}{4. stellar evolution \\& radioactive decay ages} &&  (b) Lyman-$\\alpha$ forest\\\\ \\multicolumn{2}{l}{5. distance scale}\\\\ & (a) distance ladders\\\\ & (b) gravitational lensing\\\\ \\hline \\end{tabular} \\end{table}   This situation is illustrated in Table~1 (adapted from the book, {\\it Finding the Big Bang} \\cite{FTBB}; details are there and in other reviews of the state of the cosmological tests). Each measurement named in the table is an independent (or close to it) probe of the universe. Each is capable of falsifying $\\Lambda$CDM. This cosmology passes the tests so far. The measurements are difficult, and issues of interpretation and systematic errors are well worth continued close examination. But the diversity of these  probes, and the consistency of independent constraints on the relevant parameters in $\\Lambda$CDM, make a close to compelling case that we have not been seriously misled. That certainly does not mean the $\\Lambda$CDM cosmology is the final word on the structure and physics of the universe that will be needed to interpret coming generations of measurements. But we can be sure that if $\\Lambda$CDM is replaced by something better the new theory will predict a universe that looks much  like $\\Lambda$CDM, because that is what is observed.  \\subsection{Precision Cosmology and Accurate Cosmology}  We should take another moment to admire the abundance of probes of the large-scale nature of the universe represented in Table~1, and to recall the importance of the diversity of these measurements.  It has been said that we are moving toward precision cosmology, which is correct but misses the point.  A digital scale may read out the weight of an object to many significant figures, in a precise measurement. But if the scale is not well calibrated the measurement may not be very accurate. That is, accuracy is what remains of precision after discounting systematic errors. And our goal is an accurate cosmology.  In the pursuit of precision cosmology it is natural to concentrate on the relatively few observations that lend themselves to the most precise measurements. This is good science, provided one bears in mind that if the precision measurements are not more numerous than the parameters that may be adjusted to fit them we are getting a cosmology of dubious accuracy. I offer one example, the properties of galaxies.  The mass distributions in the outer parts of galaxies were an apparent anomaly within the old cosmology, and a hint of something new: dark matter. Recent research proceeds in the opposite direction: take $\\Lambda$CDM, dark matter and all, as given, to be used as the basis for analyses of how galaxies acquired their observed  properties. Since galaxies are complicated --- our understanding of star formation and its impact on the diffuse baryons is particularly schematic --- the program must add parameters that are physically well motivated but adjustable because we don't understand how the physics actually is expressed. The freedom of adjustment makes it difficult to judge the significance of the notable advances in this program. The pursuit of more accurate cosmology may be aided by a return to the earlier direction. There are apparent anomalies in the properties of galaxies in terms of what might have been expected in a straightforward reading of numerical simulations of structure formation in $\\Lambda$CDM. Instead of asking how the parameters needed to describe the histories of the baryons may be adjusted to remedy this situation consider instead whether apparent anomalies might be hints to better fundamental physics. It is not difficult to invent alternative physics that fits the tests on the scale of the Hubble length about as well as $\\Lambda$CDM, but on the scale of galaxies does something different, interesting, and just possibly better. An example of an alternative theory on the scale of galaxies is discussed next; others are in \\cite{Khoury} -- \\cite{Carroll et al.}. ",
        "conclusions": "Cosmology has changed from the mode of operation the older of us remember without nostalgia. We are now tightly constrained by a rich observational basis that makes it exceedingly delicate --- though certainly not impossible --- to find and demonstrate viability of alternatives to the $\\Lambda$CDM cosmology. We each have chosen a mode of research appropriate for this relatively new situation in cosmology. It is worth listing strategies, to remind ourselves and funding agencies that we do have choices.  It is entirely reasonable to take the conservative position that $\\Lambda$CDM, maybe with slow evolution of the dark energy density and fine tuning of initial conditions, includes all the physics that is going to be relevant to analyses of phenomenological cosmology. Then the goal becomes to understand how this given physics is expressed in the processes that make the universe we observe, from the Hubble length down to galaxies and stars and planets, and back in time to light element production and maybe baryogenesis. This strategy has proven its worth by defining a host of productive research problems.  If $\\Lambda$CDM does differ from reality enough to matter it means this conservative strategy will reveal it  through identification of phenomena that cannot be reconciled with the cosmology. But beware of precision cosmology, because the observations that lend themselves precision measurements reduce the tests in Table~1 to a number perilously close to the parameters we are willing to adjust in $\\Lambda$CDM. (Currently discussed additions to the more traditional cosmological parameters include evolution and maybe spatial gradients of the dark energy density, a soup\\c{c}on of annihilating or collisional or warm dark matter, or maybe  even a new long-range interaction in the dark sector. A competent theorist likely could offer quite a few other ways to save the phenomena within $\\Lambda$CDM.) I expect that the better strategy for finding a more accurate cosmology will be to examine all measurements, however inaccurate, that probe different aspects of the universe and may challenge standard ideas.  A more proactive strategy for discovering a possibly more accurate cosmology is to search for apparent anomalies in $\\Lambda$CDM, situations that seem particularly challenging for this theory. My list commences with a pronounced difference between the physics of the visible and invisible universes. Physics in the visible sector is wonderfully simple, in the appropriate sense, but its expression is spectacularly complicated. Physics in the invisible sector is so exceedingly simple that its expression is simple: the dark matter just piles up in halos and the dark energy sits there or maybe slowly evolves. It is wishful thinking, but also sensible, to suspect that this apparently anomalous situation is telling us that we have settled for the simplest physics we can get away with at the present crude levels of probes of this sector. It presents no hints of phenomena to look for; I read it an invitation to look broadly.  Other apparent anomalies may offer more direct hints of better physics. Examples include preliminary results from underground dark matter particle detectors and the search for signatures of dark matter annihilation. At the time of writing the experimental/observational situation is confused, maybe in part a result of systematic errors in difficult measurements, maybe in part because we are seeing hints to something new. Other examples are the more curious properties of galaxies.  How did galaxies arrive at their scaling relations that show such little regard for environment? Why do the voids defined by normal galaxies look so very empty? Why do most of the nearest spiral galaxies edge up to the very empty Local Void? How did massive black holes in active galactic nuclei form so early in the expansion of the universe? Why is large-scale structure in the galaxy distribution so very large?  If such curiosities are expressions of standard physics then understanding how they came about will teach us something of value. And maybe some are hints to better physics, which would be a good thing too.  Research in the proactive direction may accept the standard physics of the visible universe and seek something more interesting in the poorly explored invisible part. As one sees in these Proceedings, there is no shortage of ideas here. The adoption of standard visible sector physics is conservative, but maybe overly so. For example, the precision tests of general relativity are on scales some fifteen orders of magnitude smaller than the Hubble length. It is sensible to continue to question this enormous extrapolation, to consider modifications of gravity physics. This strategy has led to viable alternatives that to date seem to me to be less logically compelling than general relativity with $\\Lambda$CDM. This is not a very convincing argument, but I expect it will lead most in the community to continue to work within the physics of $\\Lambda$CDM unless/until some new phenomenology or really attractive idea drives us from it.  I refer to Wigner's \\cite{Wigner} essay, {\\it The Unreasonable Effectiveness of Mathematics in the Natural Sciences},  for the arguments that lead me to expect that the community will not arrive at the one true cosmology, if such a thing exists. We will instead end up with the best approximation allowed by the limited ability to observe. We cannot rule out the possibility that something will turn up that invalidates our present cosmology, but from all experience that is exceedingly unlikely. Much easier to imagine is fine tuning to fit new measurements will lead to a more accurate cosmology that contains elements of $\\Lambda$CDM along with new and interesting adjustments."
    },
    "0910/0910.2554_arXiv.txt": {
        "abstract": "% The evolution of galaxy merger rates and its impact on galaxy properties have been studied intensively over the last decade. It becomes clear now that various types of mergers, i.e. gas-rich (wet), gas-poor (dry), or mixed mergers, affect the merger products in different ways. The epoch when each type of merger dominates also differs. In this talk, I review the recent progress on the measurements of galaxy merger rates out to $z \\sim 3$ and the level of interaction-triggered star formation using large samples from various redshift surveys. These results provide insights to the importance of mergers in the mass assembly history of galaxies and in the evolution of galaxy properties. I also present new results in characterizing the environment of galaxy mergers, and discuss their implications in the built up of red-sequence galaxies. ",
        "introduction": "% Galaxy mergers have long been suggested to be associated with a variety of observational phenomena, including Ultra-Luminous Infrared Galaxies (ULIRGs), Quasars, post-starburst galaxies, Active Galactic Nuclei (AGN), and Submillimeter Galaxies (SMGs), etc. In certain galaxy evolution models, these objects may actually correspond to different phases during galaxy interactions \\citep{hop06}. Galaxy mergers have significant impact on galaxy morphology, kinematics, stellar masses, and star formation histories of galaxies, as well as the mass of its central black holes. Recent studies on the galaxy luminosity function (LF) and stellar masses function (SMF) reveal that the number densities and stellar mass densities of galaxies in the red sequence have been increased by a factor of 2 to 3 since $z \\sim 1$ while that of blue forming galaxies remain similar over the same period \\citep{fab07,bel07}, suggesting a migration of galaxies from the blue cloud to the red sequence. One plausible mechanism responsible for this transformation is through merging of two blue galaxies, the so-called 'wet mergers' (gas-rich mergers) followed by subsequent mergers among re-sequence galaxies, the so-called 'dry mergers' (gas-poor mergers) \\citep{fab07}. Therefore in order to better understand how the present-day massive galaxies are assembled, and how the galaxies have been evolved, we need to have improved constraints on the abundance of mergers and where they occur. There are three aspects regarding galaxy interactions that I focus on in this talk: the absolute galaxy merger rates, the influence of mergers in triggering star formation, and the environment of merging galaxies. ",
        "conclusions": "I have argued that studies of merger rates, effectiveness of the tidally-triggered star formation, and the host environment of mergers are the three keys to pin down the role of mergers in the evolution history of galaxies. Despite that our understanding of the above issues have been dramatically improved over the last few years, there remains several pieces of details missing in the whole picture. For example, under what kind of conditions will wet mergers lead to the quench of star formation and hence become K+A galaxies? What is the outcome of mixed mergers? What is the merger rate beyond redshift one? How well can we separate the contributions of minor mergers from major mergers? And finally, what is the role of AGN during galaxy encounters? Both detailed studies of individual interacting systems and statistical results based on future larger and deeper surveys of merging galaxies, together with improved numerical modeling that incorporate more sophisticated physics will provide essential knowledge to resolve the aforementioned issues."
    },
    "0910/0910.3691_arXiv.txt": {
        "abstract": "While feedback is important in theoretical models, we do not really know if it works in reality. Feedback from jets appears to be sufficient to keep the cooling flows in clusters from cooling too much and it may be sufficient to regulate black hole growth in dominant cluster galaxies. Only about 10\\% of all quasars, however, have powerful radio jets, so jet-related feedback cannot be generic.  The outflows could potentially be a more common form of AGN feedback, but measuring mass and energy outflow rates is a challenging task, the main unknown being the location and geometry of the absorbing medium. Using a novel technique, we made first such measurement in NGC 4051 using XMM data and found the mass and energy outflow rates to be 4 to 5 orders of magnitude {\\it below} those required for efficient feedback. To test whether the outflow velocity in NGC 4051 is unusually low, we compared the ratio of outflow velocity to escape velocity in a sample of AGNs and found it to be generally less than one. It is thus possible that in most Seyferts the feedback is not sufficient and may not be necessary. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.1007_arXiv.txt": {
        "abstract": "{} % {The co-evolution of host galaxies and the active black holes which reside in their centre is one of the most important topics in modern observational cosmology. Here we present a study of the properties of obscured Active Galactic Nuclei (AGN) detected in the CDFS 1Ms observation and their host galaxies.} {We limited the analysis to the MUSIC area, for which deep K-band observations obtained with ISAAC@VLT are available, ensuring accurate identifications of the counterparts of the X--ray sources as well as reliable determination of photometric redshifts and galaxy parameters, such as stellar masses and star formation rates. In particular, we: 1) refined the X-ray/infrared/optical association of 179 sources in the MUSIC area detected in the Chandra observation; 2) studied the host galaxies observed and rest frame colors and properties.} {We found that  X--ray selected (L$_X\\gs10^{42}$ erg s$^{-1}$) AGN show Spitzer colors consistent with both AGN and starburst dominated infrared continuum; the latter would not have been selected as AGN from infrared diagnostics. The host galaxies of X--ray selected obscured AGN are all massive (M$_*>10^{10}$ M$\\odot$) and, in 50\\% of the cases, are also actively forming stars (1/SSFR$<$ t$_{Hubble}$) in dusty environments. The median L/LEdd value of the active nucleus is between 2\\% and 10\\% depending on the assumed M$_{BH}$/M$_{*}$ ratio. % Finally, we found that the X--ray selected AGN fraction increases with the stellar mass up to a value of $\\sim30$\\% at z$>1$ and M$_*>3\\times10^{11}$ M$\\odot$, a fraction significantly higher than in the local Universe for AGN of similar luminosities.} {} ",
        "introduction": "Galaxy interactions, and more in general the large scale structure (LSS) galaxy environment, are thought to play a major role in regulating both star-formation and accretion onto nuclear super-massive black holes (SMBHs). This implies that the full understanding of galaxy evolution requires a good knowledge of the SMBH census through cosmic time. In particular, feedback between the central SMBHs in their active phases and the interstellar medium is likely to affect strongly the evolution of their host galaxies.   A short, powerful but highly obscured growth phase of both SMBHs and their host galaxies is predicted by many models for the co-evolution of galaxies and AGN (Silk \\& Rees 1998, Fabian 1999, Granato et al. 2004, Di Matteo et al. 2005, Menci et al. 2008). This phase ends when strong AGN winds and shocks heat the interstellar medium, blowing away the dust and gas and inhibiting further star-formation in the AGN host galaxies. According to this ``evolutionary sequence'' (e.g. Hopkins et al. 2008), highly obscured AGN should be associated to young galaxies in the process of assembling most of their stellar mass through significant episodes of star formation. On the contrary, unobscured AGN should be associated to galaxies with low or absent episodes of star-formation, given that most of the gas and dust responsible for the star formation has been blown away by the effect of AGN feedback. Many observational evidences (see Alexander et al. 2005, Page et al. 2004, Stevens et al. 2006) and theoretical arguments (Menci et al. 2008 and references therein) in favor of the evolutionary sequence do exist. These results challenge our 20-years old AGN view, in which the differences we see in different classes of sources - especially between \"obscured\" and \"unobscured\" ones - are simply due to orientations effects (Antonucci et al. 1985, Antonucci 2003, Urry \\& Padovani 1995). \\\\ A correct and complete identification of obscured and highly obscured AGN at high redshift is therefore crucial because they may represent the first, still little explored, phase of the common growth of both SMBHs and their host galaxies.  Moderately and highly obscured AGN at z$\\sim2$, where most of the accretion and star-formation processes are on-going, offer an ideal tool for a direct test of feedback models. In these objects the nuclear light is completely blocked or strongly reduced and does not overshine the galaxy optical and near-infrared light. This gives us the possibility to study host galaxy morphologies, colors and spectral energy distributions without the difficulty of disentangling star-light from nuclear light. We can therefore study galaxy properties during active phases, e.g. {\\it when nuclear feedback should be in action}.  Obscured AGN at cosmological distances are however usually faint in the observed optical bands, because the UV rest frame is strongly reduced by dust extinction. On the other hand, moderately obscured AGN (or Compton thin, $10^{22}<$N$_H<10^{24}$ cm$^{-2}$) are common in X-ray images (Bauer et al. 2004, Comastri \\& Brusa 2008 and reference therein), making up $\\sim$50\\% of the full X-ray population at fluxes $<10^{-14}$ \\cgs (Gilli, Comastri \\& Hasinger 2007) once the contribution from the normal starforming galaxy population is removed (Ranalli et al. 2005). Compton thick AGN (CT, N$_H\\gs10^{24}$ cm$^{-2}$) are faint also in the X-ray band, because photoelectric absorption and Compton scattering strongly reduce the X-ray flux up to $>10$ keV, and over the entire X--ray range if N$_H\\gs10^{25}$ cm$^{-2}$. Indeed, only a dozen CT AGN is present in the deepest images of the X-ray sky, the Chandra Deep Fields (Norman et al. 2002, Tozzi et al.  2006, Georgantopoulos et al. 2008).  Identification of the correct counterparts of obscured AGN at cosmological redshift is not trivial. These objects are faint in the optical band, because 1) the intrinsic AGN emission is absorbed by the surrounding material, and 2) the host galaxy light is strongly reduced by cosmological dimming (L$^*$galaxies at z=1-2 have R$\\sim24-25$). The probability to find by chance a galaxy of these magnitudes in the Chandra error boxes is not negligible.  The identification process is made easier by using deep near infrared images, because the surface density of these sources is smaller than that of optical ones, bringing the probability of finding a galaxy by chance in the Chandra error boxes to comfortably small values.  Moreover, the K-band flux is more tightly correlated with the X-ray flux than the optical (obscured) one.  For this reason, we have reanalyzed the identifications of the Chandra Deep Field South (CDFS) 1 Ms sources (Alexander et al. 2003) in the area covered by sensitive K band and IRAC 3.6$\\mu$m and 4.5$\\mu$m observations (the GOODS MUSIC area, Grazian et al. 2006).  The excellent multiband photometry available in this area allows the determination of a reliable photometric redshift for the faint sources not reachable by optical spectroscopy.  Once identified the counterparts of the X-ray sources and their redshift, it is possible to proceed to a detailed study of the properties of their host galaxies. We present here a study of the properties of the host galaxies of X-ray obscured AGN, such as their morphology and close environment, optical and infrared rest frame colors, mid infrared to optical spectral energy distributions. Galaxy masses and star-formation rates are derived from the fit of galaxy templates to the optical to near infrared spectral energy distributions. We finally compare the masses and star-formation rates of obscured AGN host galaxies to those of inactive galaxies in the field selected in both optical and infrared bands.  The paper is organized as follows: Sect. 2 presents the CDFS datasets, the X-ray to optical/infrared association and the obscured AGN sample. Section 3 presents the observed frame colors of the obscured AGN sample.  Section 4 presents the host galaxies properties (masses and star-formation rates, SFRs ) of the sample and our estimates of the fraction of AGN in mass selected samples.  Section 5 presents a a discussion of the results and Section 6 outlines a summary of the most important points.  A cosmology with $H_0=70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_M$=0.3, $\\Omega_{\\Lambda}=0.7$ is adopted throughout (Spergel et al. 2003). ",
        "conclusions": "We presented a new catalog of the counterparts of the 179 extragalactic X--ray sources detected in the 1Ms Chandra observation of the MUSIC/CDFS/GOODS field and an extensive analysis of the host galaxies properties of obscured AGN. \\\\ We quantified the bias in the determination of the counterparts of X--ray selected sources when the match is limited to optical catalogs only, with respect to the combined use of optical and near infrared (deep K--band and IRAC) data.  We estimate that the fraction of misidentified X--ray sources previously reported in the literature is of the order of $\\sim 6\\%$, and rises up to $\\sim14$\\% when optically faint ($z>24$) sources are considered (see Figure 1); the use of an optically-based catalog biases the identification against the most extreme, obscured sources therefore preventing the exact knowledge of the multiwavelength properties of the X-ray counterparts. \\\\ \\par\\noindent In order to study the host-galaxies of obscured AGN, we defined a sample of 116 ``bona fide'' obscured AGN, by selecting sources without broad lines in the optical spectra and with small optical nuclear emission with respect to the host galaxy optical emission (see Figure 3). From eyeball inspection of the host galaxy morphology, we found a variety of cases, and a disturbed morphology (due to activity/merging/star formation) in more than half of the sub-sample for which a morphological classification could be made (see Figure 4).\\\\ \\par\\noindent We investigated the optical to infrared colors of these obscured AGN. The most striking result is that half of the X--ray selected obscured AGN in the redshift range z=1.5-4.0 have a 8.0$\\mu$m to 4.5$\\mu$m flux ratio $<2$ and according to Pope et al. (2008) would have been classified as ``star-formation'' dominated objects (see Figure 6); 50\\% of them in addition have 24.0$\\mu$m to 8.0$\\mu$m flux ratio $<5$, where the original Pope et al. (2008) diagram is almost empty. Moreover, previous analysis based on Chandra stacking analysis of sources with high MIR/O (Daddi et al. 2007a, Fiore et al. 2008, 2009) claimed a large contribution ($\\geq80$\\%) of heavily obscured (Compton Thick) sources among the stacked population (but see also discussion in Donley et al. 2008). This suggests that 1) the accretion activity in high-redshift sources is unambiguously revealed thanks to the presence of a strong X--ray emission, and 2) the star-formation region as defined by Pope et al. (2008) contains not only objects in which the bolometric luminosity is dominated by the star-formation processes, but also a not-negligible number of objects hosting candidate obscured/Compton Thick  AGN. In conclusion, our obscured AGN show Spitzer colors consistent with both an AGN dominated continuum and a starburst dominated continuum in the MIR. \\\\ \\par\\noindent We found that the hosts of obscured AGN are redder in (U-V) rest frame than the overall galaxy population at the same redshift: in particular, obscured AGN mainly populate the red sequence and the green valley in the color-magnitude plots (Figure 7, left panel), in agreement with the results of Nandra et al. (2007), and Silverman et al. (2007). For the MUSIC sample the U-V galaxy colors are strongly correlated with the K band absolute magnitude (Figure 8, left panel), and therefore with the galaxy stellar mass, with the most massive systems having a redder color.  The hosts of the obscured AGN are therefore found in the red-massive tail of the distribution of optically selected galaxies in all three redshift bins considered. AGN feedback is often invoked as one of the main responsible for the observed galaxy colors (Nandra et al. 2007, Hasinger 2008). However, it is well known that the main ingredient for nuclear activity is the presence of a SMBH in galaxy nuclei, and that SMBHs are found nearly exclusively in massive galaxies (e.g. Magorrian et al. 1998). Therefore, it is not truly surprising to find AGN hosted in massive galaxies and the simple presence of AGN in massive red galaxies is not enough to argue for a significant feedback effect on the observed colors, because of the strong color-mass correlation. Were AGN feedback responsible for the observed red colors, since galaxy colors are strongly correlated with the galaxy mass and AGN are found preferably in massive galaxies, then AGN feedback should be also considered as one of the main players in the building of the galaxy mass-color correlation. \\\\ \\par\\noindent We found that about 2/3 of the obscured AGN hosts at all redshifts show substantial ($>10$ M$_\\odot$ yr$^{-1}$) star formation activity (Figure 9 left panel) and about half live in galaxies which are still actively forming stars with respect to their mass (Figure 9, right panel). For these sources, the observed red colors are likely due to dust extinction rather than evolved stellar population. We then conclude that a significant fraction of obscured AGN live in massive, dusty star-forming galaxies with red optical colors. This result is in qualitative agreement with the morphological analysis. Higher luminosity X-ray selected AGN are not systematically found in objects with the highest SSFR (see Fig. 9 right panel), in agreement with Alonso-Herrero et al. (2008). \\\\ \\par\\noindent We compared the number of obscured AGN and of all X--ray selected AGN to the number of field galaxies in broad bins of galaxy stellar mass (M$_*=10^{10}-10^{12}$ M$_\\odot$) and redshifts (z=0.6-1, z=1-2, z=2-4). We find that the AGN fraction increases with the host galaxy stellar mass, from $\\sim$1\\% at M$_*\\sim 10^{10}$ M$_\\odot$ to $\\sim$30\\% at M$_*\\sim 3\\times10^{11}$ M$_\\odot$ (see also Yamada et al. 2009), and the actual trend of increasing AGN fraction  as a function of the stellar mass is probably steeper given the uncompleteness of the MUSIC sample at M$_{*}<10^{11}$ M$\\odot$ (see Section 4.4). The uncertainties on the stellar mass estimate from the SED fitting have the effect of shifting at lower masses (of $\\sim 0.25$ dex) the datapoints, leaving the total fraction and the slope unchanged. \\\\ We compared this trend with that observed in the local Universe (Best et al. 2005) for AGN with luminosity above similar thresholds. While the observed trend is the same, in all the investigated redshift bins the AGN fraction is higher than that observed in the local Universe, and it could likely be even higher. In fact, we are comparing here AGN selected with two different methods: forbidden emission line luminosity (SDSS) and X-ray emission (GOODS). The latter sample does not contain most Compton thick AGN. On the other hand, Compton thick AGN may well be present in [OIII] selected AGN samples. The fraction of Compton thick AGN not directly detected in deep Chandra surveys is estimated between 40\\% and 100\\% of the X-ray selected AGN, using infrared selection or other techniques (see Donley et al. 2008, Fiore et al. 2009 and references therein). Therefore, under the simplest assumption that this Compton Thick AGN fraction is constant with the galaxy mass, the discrepancy observed in Figure 10 can increase by up of a factor of 2. \\\\ The fraction of active galaxies to the total galaxy population is proportional to the AGN duty cycle. Our results would thus suggest higher AGN duty cycles at z=2-4 than at z=0, in agreement with expectations from most recent semi-analytic models (e.g. Menci et al. 2008), in which at higher redshift the AGN activity is present in a larger number of galaxies than locally.  \\\\ \\par\\noindent The fact that the most luminous obscured AGN are found in the most massive galaxies at all investigated redshifts may suggest that the L/L$_{Edd}$ of the obscured AGN is similar, particularly in the case of the most luminous sources (logLx$>43$ erg s$^{-1}$), for which the threshold in luminosity introduces a bias against the sources accreting at lower rates in the lowest redshift bin. Assuming the local Magorrian relation between  M$_{BH}$ and M$_*$ (e.g Marconi \\& Hunt 2003) and a bolometric correction of 20 (e.g. Marconi et al. 2004) the median values of L$_X$/M$_*$=32.83 and L$_X$/M$_*$=32.87 correspond to L/L$_{Edd}\\sim0.1$. Although suffering from large uncertainties associated with the stellar mass and BH mass estimates (see Section 4.4), this value can be considered as representative of the accretion state of the most luminous, obscured AGN in the present sample. Similar results are obtained from a comprehensive analysis of host galaxies properties of Chandra Deep Field North X--ray sources sources at z=2-4 (Yamada et al. 2009) and are also typical of unobscured Type 1 AGN at z$>1$ (Merloni et al. 2010, Trump et al. 2009). \\\\ It is instructive to compare these findings with similar estimates in the local Universe.  Kauffmann \\& Heckman (2009) using SDSS data found an average L/L$_{Edd}$ value of $\\sim 0.01$ and a log normal distribution for this parameter, for AGN hosted in galaxies with significant on-going star formation, while AGN in inactive host galaxies follow a power-law distribution. As discussed in section 4.4, and shown in the upper right panel of figure 11, we also find marginal evidence of different accretion rates distributions for the populations of AGN in 'active' and 'inactive' host galaxies. \\\\"
    },
    "0910/0910.3833_arXiv.txt": {
        "abstract": "In order to investigate whether galaxy structures are compatible with the predictions of the standard LCDM cosmology, we focus here on the analysis of several simple and basic statistical properties of the galaxy density field. Namely, we test whether, on large enough scales (i.e., $r>10$ Mpc/h), this is self-averaging, uniform and characterized by a Gaussian probability density function of fluctuations. These are three different and clear predictions of the LCDM cosmology which are fulfilled in mock galaxy catalogs generated from cosmological N-body simulations representing this model.  We consider some simple statistical measurements able to tests these properties in a finite sample. We discuss that the analysis of several samples of the Two Degree Field Galaxy Redshift Survey and of the Sloan Digital Sky Survey show that galaxy structures are non self-averaging and inhomogeneous on scales of $\\sim$ 100 Mpc/h, and are thus intrinsically different from LCDM model predictions. Correspondingly the probability density function of fluctuations shows a \"fat tail\" and it is thus different from the Gaussian prediction. Finally we discuss other recent observations which are odds with LCDM predictions and which are, at least theoretically, compatible with the highly inhomogeneous nature of galaxy distribution. We point out that inhomogeneous structures can be fully compatible with statistical isotropy and homogeneity, and thus with a relaxed version of the Cosmological Principle. ",
        "introduction": "In the past twenty years many observations have been dedicated to the study of the large scale distributions of galaxies \\cite{cfa1,cfa2,pp,ssrs2,lcrs,sdss,colless01}.  In particular during the last decade two ambitious observational programs have measured the redshift of more than one million objects \\cite{sdss,colless01}.  All these surveys have detected larger and larger structures, thus finding that galaxies are organized in a complex network of clusters, super-clusters, filaments and voids. For instance the famous ``slice of the Universe'', that represented the first set of observations done for the CfA Redshift Survey in 1985 \\cite{dlhg85}, mapped spectroscopic observations of about 1100 galaxies in a strip on the sky 6 degrees wide and about 130 degrees long. This initial map was quite surprising, showing that the distribution of galaxies in space was anything but random, with galaxies actually appearing to be distributed on surfaces, almost bubble like, surrounding large empty regions, or ``voids.''.  The structure running all the way across the survey between 50 and 100 Mpc/h\\footnote{We use $H_0=100 h$ km/sec/Mpc for the value of the Hubble constant.} was called the ``Great Wall'' and at the time of the discovery was the largest single structure detected in any redshift survey. Its dimensions, limited only by the sample size, are about $200\\times 80 \\times 10 $ Mpc/h, a sort of like a giant quilt of galaxies across the sky \\cite{gh89}. More and more galaxy large scale structures were identified in the other redshift surveys such as the Perseus-Pisces super-cluster \\cite{pp} which is one of two dominant concentrations of galaxies in the nearby universe. This long chain of galaxies lies next to the the so-called Taurus void, which is a large circular void bounded by walls of galaxies on either side of it. The void has a diameter of about 30 Mpc/h.  Few years ago, in the larger sample provided by the Sloan Digital Sky Survey (SDSS), it has been discovered the Sloan Great Wall \\cite{sloangreatwall}, which is a giant wall of galaxies which may be the largest known structure in the Universe, being nearly three times longer than the Great Wall. The discovery of larger and larger structures was surprising because standard cosmological models unambiguously predict that fluctuations should be small on large scales with rapidly decaying correlations; i.e., the spatial extension of structures in these models should be limited to some tens Mpc/h. ",
        "conclusions": "We interpret the systematic differences found in the behavior of the PDF of conditional fluctuations as due to a systematic effect in the fact that sample volumes are not large enough for conditional fluctuations, filtered at such large scales, to be self-averaging, i.e.  to contain enough structures and voids of large size to allow a reliable determination of average (conditional) quantities.  We pointed out the problems related to the estimation of amplitude of fluctuations and correlation properties from statistical quantities which employ the normalization to the estimation of the sample average. As long as a distribution inside the given sample is not self-averaging, and thus not homogeneous, the estimation of the two-point correlation function is necessarily biased by strong finite size effects. Our results are incompatible with homogeneity at scales smaller than $\\sim 100$ Mpc/h. While these results are at odds with LCDM predictions, they can be compatible, at least theoretically, with a several recent observations which also pose fundamental challenges to such a model.   For instance, Kashlinsky et al. \\cite{Kashlinsky}, by studying the fluctuations in the cosmic microwave background generated by the scattering of the microwave photons by the hot X-ray emitting gas inside clusters, have measured a coherent flow out to 300 Mpc/h with a fairly high amplitude of 600-$10^3$ km/sec.  This is incompatible with the standard LCDM model predictions. Indeed, on such large scales the theoretical predictions are very simple, because in these models gravitational clustering is still linear at those scales as density fluctuations are small on large scales.  Similarly Watkins et al. \\cite{Watkins} estimated the bulk flow in all major peculiar velocity surveys finding that the data suggest that the bulk flow within a Gaussian window of radius 50 Mpc/h is 407 km/s.  They noticed that this large-scale bulk motion indicates that there are significant density fluctuations on very large scales. Indeed, a flow of this amplitude on such a large scale is not expected in the LCDM model cosmology, for which the predicted one-dimensional r.m.s. velocity is about 110 km/s.  Thus the same problem for the model predictions we found in the galaxy density field, are found also for the galaxy velocity field. They may have the same origin, namely the fact that there are galaxy structures which are too extended in space and have a too large amplitude to be compatible with the predictions of the standard model. In the case of the velocity field there is another important element: the velocity field is generated by all the mass and not only by the luminous component. Thus if the velocity field is so high on such large scales, this may imply that this is generated by the large scale inhomogeneities present in the overall mass distribution, i.e. luminous plus dark. The fact that the whole mass distribution is not homogeneous is compatible also with the results on the matter density field derived by gravitational lens observations, where very extended structures have been found \\cite{massey}.  Thus an important point which we aim to investigate in future works, concerns the characterizing of the gravitational field generated by galaxies.  When a distribution is inhomogeneous important contributions to the gravitational force acting on a point can be due to faraway sources \\cite{force}.  The relation between the large scale inhomogeneities of galaxy distribution and other observational data should be examined in detail. For instance a statistically significant anisotropy of the Hubble diagram at redshifts $z < 0.2$ was discovered by \\cite{sw08}. A local violation of statistical isotropy and homogeneity, which may very well happen when matter distribution is inhomogeneous \\cite{sl94,pwa}, can be related to such findings. although it is not excluded that a systematic error in the observations or data analysis affect these results.  Finally, it is worth mentioning that Joyce et al. \\cite{pwa} pointed out that an inhomogeneous distribution, as a fractal, does not preclude the description of its gravitational dynamics in the framework of the Friedmann-Robertson-Walker solutions to general relativity. Indeed, the problem is often stated as being due to the incompatibility of a fractal with the Cosmological Principle, where this principle is identified with the requirement that the matter distribution be isotropic and homogeneous. This identification is in fact very misleading for a non-analytic and inhomogeneous structure like a fractal, in which all points are equivalent statistically, satisfying what has been called a Conditional Cosmological Principle \\cite{sl94,book}.  By treating the fractal as a perturbation to an open cosmology in which the leading homogeneous component is the cosmic background radiation one may get a simple explanation for the supernovae data which indicate the absence of deceleration in the expansion. This is indeed a very simplified theoretical model to interpret the SN data and the large scale inhomogeneities in framework of the Friedmann-Robertson-Walker metric.    \\begin{theacknowledgments} I am in debt with Yuri V. Baryshev and Nickolay L. Vasilyev for many collaborations on this topic. I also thank Tibor Antal, Michael Joyce, Andrea Gabrielli, Martin Lopez-Correidoira and Luciano Pietronero for useful remarks and discussions. I am grateful to Michael Blanton and David Hogg and for interesting comments. \\end{theacknowledgments}"
    },
    "0910/0910.1836_arXiv.txt": {
        "abstract": "We present optical spectroscopic follow-up of a sample of Distant Red Galaxies (DRGs) with $K^{tot}_{s,Vega}<22.5$, selected by $(J-K)_{Vega}>2.3$, in the Hubble Deep Field South (HDFS), the \\1054 field, and the Chandra Deep Field South (CDFS).  Spectroscopic redshifts were obtained for 15 DRGs.  Only 2 out of 15 DRGs are located at $z<2$, suggesting a high efficiency to select high-redshift sources.  From other spectroscopic surveys in the CDFS targeting intermediate to high redshift populations selected with different criteria, we find spectroscopic redshifts for a further 30 DRGs.  We use the sample of spectroscopically confirmed DRGs to establish the high quality (scatter in $\\Delta z/(1+z)$ of $\\sim 0.05$) of their photometric redshifts in the considered deep fields, as derived with EAZY (Brammer et al. 2008).  Combining the spectroscopic and photometric redshifts, we find that 74\\% of DRGs with $K^{tot}_{s,Vega} < 22.5$ lie at $z>2$.  The combined spectroscopic and photometric sample is used to analyze the distinct intrinsic and observed properties of DRGs at $z<2$ and $z>2$.  In our photometric sample to $K^{tot}_{s,Vega} < 22.5$, low-redshift DRGs are brighter in $K_s$ than high-redshift DRGs by 0.7 mag, and more extincted by 1.2 mag in $A_V$.  Our analysis shows that the DRG criterion selects galaxies with different properties at different redshifts.  Such biases can be largely avoided by selecting galaxies based on their rest-frame properties, which requires very good multi-band photometry and high quality photometric redshifts. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.0001_arXiv.txt": {
        "abstract": "We use the ultra-deep WFC3/IR data over the HUDF and the Early Release Science WFC3/IR data over the CDF-South GOODS field to quantify the broadband spectral properties of candidate star-forming galaxies at $z\\sim7$.  We determine the $UV$-continuum slope $\\beta$ in these galaxies, and compare the slopes with galaxies at later times to measure the evolution in $\\beta$.  For luminous $L_{z=3}^{*}$ galaxies, we measure a mean $UV$-continuum slope $\\beta$ of $-2.0\\pm0.2$, which is comparable to the $\\beta\\sim-2$ derived at similar luminosities at $z\\sim5-6$.  However, for the lower luminosity $0.1L_{z=3}^{*}$ galaxies, we measure a mean $\\beta$ of $-3.0\\pm0.2$. This is substantially bluer than is found for similar luminosity galaxies at $z$$\\sim$4, just 800 Myr later, and even at $z$$\\sim$5-6. In principle, the observed $\\beta$ of $-3.0$ can be matched by a very young, dust-free stellar population, but when nebular emission is included the expected $\\beta$ becomes $\\geq-2.7$.  To produce these very blue $\\beta$'s (i.e., $\\beta\\sim-$3), extremely low metallicities and mechanisms to reduce the red nebular emission seem to be required. For example, a large escape fraction (i.e., $f_{esc}\\gtrsim0.3$) could minimize the contribution from this red nebular emission.  If this is correct and the escape fraction in faint $z\\sim7$ galaxies is $\\gtrsim$0.3, it may help to explain how galaxies reionize the universe. ",
        "introduction": "The spectral properties of high-redshift galaxies must undergo dramatic changes at some point in the past, as the metallicities in these systems drop to lower values and these systems become progressively younger.  In the limit of low metallicities, gas is no longer able to cool efficiently, likely resulting in massive extremely low-metallicity (or Population III) stars whose hot atmospheres are expected to result in a very hard $UV$-continuum spectrum and strong HeII 1640 emission.  However, the strong UV flux coming from hot stars is expected to be largely offset by the redder nebular continuum light produced by the ionized gas surrounding these massive stars (e.g., Schaerer 2002).  The newly installed WFC3/IR camera on the Hubble Space Telescope permits us to observe faint $z\\gtrsim7$ galaxies $\\gtrsim40\\times$ more efficiently than before, providing us with our most detailed look yet at the $UV$ light and spectral properties of $z\\gtrsim7$ galaxies. Already $\\gtrsim$25 likely $z$$\\sim$7-8 galaxies have been identified in the early WFC3/IR data over the Hubble Ultra Deep Field (HUDF: Oesch et al.\\ 2009b; Bouwens et al.\\ 2009b; McLure et al.\\ 2009; Bunker et al.\\ 2009), and $\\gtrsim$20 $z\\sim7$ candidate galaxies in the WFC3/IR Early Release Science (ERS) observations over the CDF-South (R.J. Bouwens et al.\\ 2009, in prep).  Here we take advantage of these early WFC3/IR observations to study the spectral properties of candidate $z\\sim7$ galaxies.  Our principal focus will be on the slope of $z\\sim7$ galaxy spectra in the $UV$-continuum -- since this slope is the primary observable we can derive from the available broadband imaging data with WFC3.  The $UV$-continuum slope $\\beta$ ($f_{\\lambda}\\propto \\lambda^{\\beta}$: e.g., Meurer et al.\\ 1999) provides us with a powerful constraint on the age, metallicity, and dust content of high-redshift galaxies; it has already been the subject of much study at $z\\sim3-6$ (Lehnert \\& Bremer 2003; Stanway et al.\\ 2005; Yan et al.\\ 2005; Bouwens et al.\\ 2006; Hathi et al.\\ 2008; see Bouwens et al. 2009a for a systematic study at $z\\sim2-6$) and even at $z\\sim7$ (Gonzalez et al.\\ 2009) using NICMOS data.  Throughout this work, we quote results in terms of the luminosity $L_{z=3}^{*}$ Steidel et al.\\ (1999) derived at $z\\sim3$: $M_{1700,AB}=-21.07$.  We refer to the HST F606W, F775W, F850LP, F105W, F125W, and F160W bands as $V_{606}$, $i_{775}$, $z_{850}$, $Y_{105}$, $J_{125}$, and $H_{160}$, respectively, for simplicity. Where necessary, we assume $\\Omega_0 = 0.3$, $\\Omega_{\\Lambda} = 0.7$, $H_0 = 70\\,\\textrm{km/s/Mpc}$.  All magnitudes are in the AB system (Oke \\& Gunn 1983). ",
        "conclusions": "Before discussing the extraordinarily blue $UV$-continuum slopes $\\beta\\sim-3$ found for lower luminosity galaxies at $z\\sim7$, we first consider the relatively luminous $L_{z=3}^{*}$ galaxy candidates.  These sources have measured $\\beta$'s of $-2$, which is similar to that observed for luminous galaxies at $z\\sim3-6$ (Figure~\\ref{fig:mcolor}).  Such slopes can be fit by a moderately young, subsolar ($0.2\\,Z_{\\odot}$) stellar population, with a maximum $E(B-V)$ of $\\sim$0.05 (Calzetti et al.\\ 2000 extinction law: see also Bouwens et al.\\ 2009a).  The lower luminosity ($0.1L_{z=3}^{*}$) galaxy candidates at $z\\sim7$, by contrast, have observed $\\beta$'s of $-3.0\\pm0.2$.  This is much bluer than for luminous $z\\sim7$ galaxies (by $\\Delta \\beta\\sim1$) and also much bluer than is found at $z\\sim5-6$ (Figures~\\ref{fig:colmag}-\\ref{fig:mcolor}).  This makes these galaxies of great interest since they are likely to be even younger, more metal poor, and dust-free than any galaxies known.  This is not a selection effect (\\S3.4), and cannot be attributed to Ly$\\alpha$ emission contributing to the broadband fluxes.  Ly$\\alpha$ does not move into the $J_{125}$-band (used to estimate $\\beta$) until $z\\gtrsim8.1$, and $\\lesssim10$\\% of our $z$-dropout sample extends to $z>8$.  An AGN contribution seem similarly unlikely, given the rarety of AGN signatures in faint Lyman-Break and Ly$\\alpha$-emitter samples (e.g., Nandra et al.\\ 2002; Ouchi et al.\\ 2008).  What then is the explanation for the very blue $\\beta$'s?  We explore several possibilities:  \\begin{figure} \\epsscale{1.16} \\plotone{f3.eps} \\caption{$UV$-continuum slope $\\beta$ we would expect for $z\\sim7$ galaxies as a function of age for constant star formation (CSF) models and an instantaneous burst (Schaerer 2002).  The gray band denotes the observed mean $\\beta$ and its uncertainty.  The slopes $\\beta$ derived from the stellar light (dotted lines) and the stellar + nebular light (solid lines) are shown.  While it is in principle possible to obtain $\\beta$'s of $-3$ with standard low metallicity ($\\geq$0.02 $Z_{\\odot}$) models, including the nebular emission associated with hot stars make the predicted $\\beta$'s $\\geq-2.7$ and hence too red.\\label{fig:dbeta}} \\end{figure}  \\begin{figure} \\epsscale{1.16} \\plotone{f4.eps} \\caption{$UV$-continuum slope $\\beta$ Schaerer (2003) calculated as a function of age for instantaneous burst models for different metallicities.  Both the slopes $\\beta$ derived from the stellar light (dotted lines) and the stellar + nebular light (solid lines) are shown.  The pure stellar light (dotted lines) has very blue $UV$-continuum slopes $\\beta$'s ($\\beta\\lesssim-3$) for all the low metallicity cases considered here.  Changing the IMF does not appear to change the conclusion here in any significant way.  Of course, these same hot low metallicity stars also ionize the gas around them, thus producing a substantial amount of redder nebular continuum light.  This makes the total SED of a galaxy much redder in general, and in the calculations by Schaerer (2003) shown here, $\\beta$ never becomes bluer than $-$3.0. \\label{fig:dbeta2}} \\end{figure}  \\subsection{Standard stellar population models}  A first question is whether it is possible to obtain a $\\beta$ of $\\sim-3$ using standard stellar population models (e.g., Leitherer et al.\\ 1999; Bruzual \\& Charlot 2003).  The answer is that it is possible, but only for very young ($<$5 Myr) star-forming systems (see Figure~\\ref{fig:dbeta}).  However, to do so ignores the nebular continuum emission from the ionized gas around the young stars. Including this nebular continuum emission can redden the observed $\\beta$ by as much as $\\Delta\\beta\\sim$0.5.  Figure~\\ref{fig:dbeta} also shows the $\\beta$'s predicted for several low metallicity (0.02 $Z_{\\odot}$ and 0.2 $Z_{\\odot}$) starbursts, as a function of age for the Schaerer (2002) stellar population models (which -- like \\textit{Starburst99} Leitherer et al.\\ 1999 -- include a nebular contribution).  In the best cases, the models predict $\\beta$'s as blue as $-2.7$, which is redder than what we observe in our faint samples.\\footnote{To match the wavelength baseline for $\\beta$ measured here, we added $-$0.1 to the $\\beta$'s (1300-1800\\AA baseline) tabulated in the Schaerer (2002, 2003).} This would suggest that lower metallicities are needed since the standard Leitherer et al.\\ (1999) or Bruzual \\& Charlot (2003) stellar population models do not produce blue enough colors to match those found in our lower luminosity $z\\sim7$ sample.  Of course, the significance of this result is only modest ($<1.3\\sigma$), but the similarly blue colors observed for z$\\sim$8 selections (Section 3.6) suggest that we may want to take this finding at face value.  \\subsection{Extremely low metallicities ($\\leq10^{-3}\\,\\,Z_{\\odot}$)}  In Figure~\\ref{fig:dbeta2} we present the $\\beta$'s predicted by the Schaerer (2003) stellar population models (which conveniently provide predictions at extremely low metallicities) as a function of age for instantaneous bursts assuming a metallicity of $0.5\\times10^{-3} Z_{\\odot}$, $0.5\\times10^{-5} Z_{\\odot}$, and zero (population III).\\footnote{The use of instantaneous burst models allows us to explore the most extreme cases.  Other models ($\\tau$, inverse $\\tau$, CSF) produce comparable $\\beta$'s.}  From the figure, it is clear that the nebular component contributes significantly to the total light output from $\\lesssim10^{-3}$ $Z_{\\odot}$ stellar populations.  While initially somewhat red due to the nebular contribution, the predicted $\\beta$'s for these ultra-low metallicity models become much bluer at ages $>10$ Myr, eventually reaching $\\beta$'s of $-$3.  Such metallicities and ages are not necessarily unreasonable for lower luminosity galaxies at $z\\sim7$, and therefore at least one possible explanation for the very blue $\\beta$'s in our selections is that the metallicities for galaxies in our sample may be $\\lesssim10^{-3}\\,Z_{\\odot}$.  The above explanation may explain some of the very blue galaxies in our selection, but given that $\\beta\\sim-3$ only for a limited period (10-30 Myr after a burst), it seems unlikely to be the general explanation (unless updated models revise the theoretical SEDs).  \\subsection{Top-heavy IMF}  One seemingly attractive explanation for the very blue $\\beta$'s observed is through a top-heavy IMF, since galaxy stellar populations would be weighted towards massive, blue stars.  The difficulty with this explanation is that these same massive stars are extraordinarily efficient at ionizing the gas around them -- resulting in substantial nebular emission and leaving the galaxy with a net $\\beta$ no bluer (and likely redder) than the young ($<$1 Myr) bursts shown in Figures~\\ref{fig:dbeta}-\\ref{fig:dbeta2} (see also discussion in Leitherer \\& Heckman 1995).  \\subsection{Minimizing nebular emission}  A significant obstacle to matching the very blue $\\beta$'s observed is the red nebular emission associated with hot, ionizing stars.  The nebular contribution could be reduced in a number of ways, by changes to the ionization parameter, metallicity, geometry, etc.  Assessing the impact of such changes would benefit from further detailed modelling.  However, we should emphasize that regardless of changes to the nebular contribution very low metallicity models (or very young ages) appear to be needed to match the very blue $\\beta$'s observed.  \\begin{figure} \\epsscale{1.16} \\plotone{f5.eps} \\caption{Mean $UV$-continuum slopes $\\beta$'s predicted for high-redshift galaxy samples versus the escape fraction $f_{esc}$. The predictions are made using the Schaerer (2003) stellar population models (including nebular emission) for metallicities of $<10^{-3}$ $Z_{\\odot}$.  Galaxies are assumed to be observed at some random point during their star formation histories and to experience instantaneous bursts of star formation every 10-40 Myr (the dashed, solid, and dotted lines give the average $\\beta$ for the first 10, 20, and 40 Myr, respectively, of the instantaneous burst).  For $f_{esc}=1$, all of the ionizing radiation from a galaxy escapes into the IGM and hence does not contribute to ionizing the gas within a galaxy (and hence the contribution from nebular continuum emission to the total light is minimal).  On the other hand, for $f_{esc}\\sim0$ (preferred in some simulations: e.g., Gnedin et al.\\ 2008), the ionizing radiation from the hot stars does not make it out of galaxies -- resulting in substantial nebular continuum emission (see Figure~\\ref{fig:dbeta2}) and hence a much redder $\\beta$.  Perhaps the very blue $\\beta$'s observed (\\textit{shaded gray region}) could indicate that the escape fraction is larger at $f_{esc}\\gtrsim0.3$ than what has commonly been considered? \\label{fig:escapef}} \\end{figure}  \\subsection{Changes in Escape Fraction?}  One possibility that we have explored to reduce the nebular emission contribution is if the ionizing radiation leaks directly into the IGM (e.g., due to the effect of SNe on the galaxies' ISM, possibly from a top-heavy IMF: Trenti \\& Shull 2009).  We consider such a possibility schematically in Figure~\\ref{fig:escapef}, showing the time-averaged $\\beta$'s expected for stellar populations of various metallicities as a function of the escape fraction $f_{esc}$ of ionizing photons into the IGM.  For $f_{esc}$ of unity, we simply recover the $\\beta$'s from the pure stellar SEDs and for $f_{esc}$ of zero, we recover the stellar + nebular SEDs.  For fractional $f_{esc}$, we interpolate between the two extremes.  The time-averaged $\\beta$'s used for Figure~\\ref{fig:escapef} is based upon those presented in Figure~\\ref{fig:dbeta2}.  Comparing the predicted $\\beta$'s with that observed (grey-shaded region: Figure~\\ref{fig:escapef}), we see that $f_{esc}\\gtrsim$0.3 would permit us to easily match the observations.  Such an escape fraction is signi\ufb01cantly higher than the $\\sim$10\\% frequently assumed in calculations assessing the sufficiency of galaxies to reionize the universe.  Since most of the luminosity density at $z>7$ comes from low luminosity galaxies, the estimated number of ionizing photons in the $z>7$ universe could increase by factors of $\\gtrsim$3. Such a large change could provide the needed photons to reionize the universe, providing a resolution to the current debate (e.g., Bouwens et al. 2008; Oesch et al. 2009a, 2009b; McLure et al. 2009; Bunker et al. 2009; Gonzalez et al. 2009; Ouchi et al.\\ 2009; Pawlik et al.\\ 2009)."
    },
    "0910/0910.2232_arXiv.txt": {
        "abstract": "We study the evolution of black holes (BHs) on the $M_{\\rm BH}-\\sigma$ and $M_{{\\rm BH}}-M_{\\rm bulge}$ planes as a function of time in disk galaxies undergoing mergers. We begin the simulations with the progenitor black hole masses being initially below $(\\Delta \\log M_{\\rm BH,i}\\sim -2)$, on $(\\Delta \\log M_{\\rm BH,i}\\sim 0)$ and above $(\\Delta \\log M_{\\rm BH,i}\\sim 0.5)$ the observed local relations. The final relations are rapidly established after the final coalescence of the galaxies and their BHs. Progenitors with low initial gas fractions ($f_{\\rm gas}=0.2$) starting below the relations evolve onto the relations $(\\Delta \\log M_{\\rm BH,f}\\sim -0.18)$, progenitors on the relations stay there $(\\Delta \\log M_{\\rm BH,f}\\sim 0)$ and finally progenitors above the relations evolve towards the relations, but still remaining above them $(\\Delta \\log M_{\\rm BH,f}\\sim 0.35)$. Mergers in which the progenitors have high initial gas fractions ($f_{\\rm gas}=0.8$) evolve above the relations in all cases  $(\\Delta \\log M_{\\rm BH,f}\\sim 0.5)$. We find that the initial gas fraction is the prime source of scatter in the observed relations, dominating over the scatter arising from the evolutionary stage of the merger remnants. The fact that BHs starting above the relations do not evolve onto the relations, indicates that our simulations rule out the scenario in which overmassive BHs evolve onto the relations through gas-rich mergers. By implication our simulations thus disfavor the picture in which supermassive BHs develop significantly before their parent bulges. ",
        "introduction": "Observations in recent years have revealed a strong correlation in the local Universe between the central supermassive black holes (BHs) and their host galaxies as manifested in the relation between the BH mass and the bulge velocity dispersion, $M_{\\rm BH}-\\sigma$ (e.g. \\citealp{2000ApJ...539L...9F,2002ApJ...574..740T,2009ApJ...698..198G};), the bulge stellar mass $M_{{\\rm BH}}-M_{\\rm bulge}$ (e.g. \\citealp{1998AJ....115.2285M}; \\citealp{2004ApJ...604L..89H}), the concentration of light in the galaxy (e.g. \\citealp{2001ApJ...563L..11G}) and the bulge binding energy, $M_{{\\rm BH}}-M_{\\rm bulge}\\sigma^{2}$ (e.g. \\citealp{2007ApJ...665..120A}). The evolution of these relations with redshift is still unclear with some studies finding evolution in the $M_{\\rm BH}-\\sigma$ \\citep{2006ApJ...645..900W,2008ApJ...681..925W} and $M_{{\\rm BH}}-M_{\\rm bulge}$ (\\citealp{2006ApJ...649..616P,2007ApJ...667..117T,2009arXiv0911.2988D}) relations, with the high redshift BHs being overmassive for a fixed $\\sigma$ and $M_{\\rm bulge}$ compared to the local relations, whereas other studies are consistent with no redshift evolution in the observed correlations \\citep{2009arXiv0908.0328G,2009ApJ...706L.215J}. A possible explanation for this discrepancy lies in the uncertainties in observational selection biases and in the evolution in the intrinsic scatter that is typically stronger for larger BH masses \\citep{2007ApJ...670..249L}.  The observed correlations are typically explained using theoretical models relying on some form  of self-regulated BH mass growth, in which gas is fed to the central black hole until the black hole releases sufficient energy to unbind the gas and blow it away in momentum- or pressure-driven winds (e.g. \\citealp{1998A&A...331L...1S,1999MNRAS.308L..39F,2001ApJ...554L.151B,2007ApJ...665.1038C,2009ApJ...699...89C}). The observed correlations and their evolution with redshift have been reproduced in semi-analytic models (e.g. \\citealp{2006MNRAS.365...11C,2008MNRAS.391..481S}), in self-consistent numerical simulations of both isolated galaxies and galaxy mergers (\\citealp{2005Natur.433..604D,2005MNRAS.361..776S,2006ApJ...641...90R}) as well as in galaxies simulated in a full cosmological setting (\\citealp{2007MNRAS.380..877S,2008ApJ...676...33D,2009MNRAS.398...53B}). The key assumption in these models is that the galaxies undergo a brief radiatively-efficient quasar phase triggered by gas-rich galaxy mergers (e.g. \\citealp{2002ApJ...580...73T,2006ApJS..163....1H,2008ApJS..175..356H}) during which the bulk of the BH growth is taking place and the observed correlations are established.  In a previous paper (\\citealp{2009ApJ...690..802J}, hereafter J09), we showed that the merger remnants of both equal- and unequal-mass mergers of disk and elliptical galaxies satisfy the observed $M_{\\rm BH}-\\sigma$ and $M_{{\\rm BH}}-M_{\\rm bulge}$  correlations. In this Letter, we study for the first time in detail the evolution of the BH scaling relations during disk galaxy mergers using a new sample of high resolution simulations including a self-consistent formulation for BH feedback. We seed the BHs initially with masses corresponding to locations below, on and above the observed relations. Thus, in contrast to previous studies we study here for the first time also overmassive BHs lying initially above the observed relations. The BH scaling relations are defined for merger remnants that have reached their final dynamical state and it is not obvious if the scaling relations are valid during the merging process and at what stage the galaxies evolve onto the relations.  \\begin{figure} \\centering \\includegraphics[width=7.4cm]{./Figures/f1.eps} \\caption{The relative distance between the BHs (top), the total SFR (middle), and the evolution of the total BH accretion rate (bottom) as a function of time. The star symbol indicate the time of BH coalescence and the filled circles at the bottom of the plots indicate the time at which ($M_{\\rm BH}$,$\\sigma$,$M_{\\rm bulge}$) are evaluated in Figs. \\ref{MBH_evo_new_BBH}-\\ref{MBH_evo_new_ABH}.} \\label{Dist_SFR_BH-evo} \\end{figure}   \\begin{table*} \\caption{The simulated merger sample} \\scriptsize{ \\label{sims} \\centering \\begin{tabular}{c c c c c c c c c c c c c c} \\hline\\hline  Model &  $M_{\\rm BH,i1}$\\footnote[$a$]{Initial BH mass in $ 10^{5} M_{\\odot}$} &   $M_{\\rm BH,i2}$\\footnotemark[$a$] & Mass ratio & $\\alpha$ & $f_{\\rm gas}$ & $N_{\\rm{gas,tot}}$ & $N_{\\rm{disk,tot}}$  & $N_{\\rm{bul,tot}}$  & $N_{\\rm{DM,tot}}$ &  $M_{\\rm BH,f}$\\footnote[$b$]{Final BH mass in $ 10^{5} M_{\\odot}$} & $\\sigma_{\\rm bul,f}$\\footnote[$c$]{Final stellar velocity dispersion in $\\rm km/s$} & $M_{\\rm bul,f}$\\footnote[$d$]{Final bulge mass in $10^{10} M_{\\odot}$} \\\\ \\hline 11B2BH & 1.0 &  1.0 & 1:1 & 25 & 0.2 & 120 000 & 480 000 & 200 000 & 800 000 & 459  &  183.8  &  5.07  \\\\ 31B2BH & 1.0 & 1.0 & 3:1 & 25 & 0.2 &  80 000 & 320 000 & 133 333 & 533 333 & 271 & 151.8 & 1.90 \\\\ 31B8BH & 1.0 & 1.0 & 3:1 & 25 & 0.8 &  320 000 & 80 000 & 133 333 & 533 333 &   1136 &  181.6  & 2.97  \\\\ 31B2BH1 & 1.0 & 1.0 & 3:1 & 100 & 0.2 &  80 000 & 320 000 & 133 333 & 533 333  &   82.5 & 144.7 &  1.26   \\\\ \\hline 11O2BH & 159 &  159 & 1:1 & 25 & 0.2 & 120 000 & 480 000 & 200 000 & 800 000  &  600  & 174.6 &   3.38 \\\\ 31O2BH & 159 & 36.4 & 3:1 & 25 & 0.2 &  80 000 & 320 000 & 133 333 & 533 333  & 253 & 132.0 & 1.55 \\\\ 31O8BH & 159 & 36.4 & 3:1 & 25 & 0.8 &  320 000 & 80 000 & 133 333 & 533 333 &   1698 &  155.2  &  3.02  \\\\ \\hline 11A2BH & 477 & 477 & 1:1 & 25 & 0.2 & 120 000 & 480 000 & 200 000 & 800 000 &   1200 &  164.2 &  3.35 \\\\ 31A2BH & 477 & 109 & 3:1 & 25 & 0.2 &  80 000 & 320 000 & 133 333 & 533 333   & 652 & 128.7 & 1.59 \\\\ 31A8BH & 477 & 109 & 3:1 & 25 & 0.8 &  320 000 & 80 000 & 133 333 & 533 333 &    842 & 122.3 & 1.21 \\\\ \\hline \\end{tabular} } \\end{table*}  \\begin{figure*} \\centering \\includegraphics[width=14.9cm]{./Figures/f2.eps} \\caption{The evolution of the BHs in the B-merger sample  $(\\Delta \\log M_{\\rm BH,i}\\sim -2)$ as a function of time on the $M_{\\rm BH}-\\sigma$ plane (top left panels) overplotted by lines giving the observed correlation from \\citet{2002ApJ...574..740T}. The bottom left panel gives the evolution on the  $M_{\\rm BH}-M_{\\rm bulge}$ plane overplotted by the observed correlation from \\citet{2004ApJ...604L..89H}. Each point is separated by $\\Delta t=0.2 \\ \\rm Gyr$, with filled circles indicating the primary galaxy, triangles indicating the secondary galaxy, stars showing the time of BH coalescence and filled squares the final state. The panels on the right give the logarithmic offset (Eq. \\ref{eq:offset}) from the corresponding observed relation as a function of time.} \\label{MBH_evo_new_BBH} \\end{figure*} ",
        "conclusions": "In this Letter, we have studied the evolution of galaxy mergers on the $M_{\\rm BH}-\\sigma$ and $M_{{\\rm BH}}-M_{\\rm bulge}$ planes as a function of time. We have shown that progenitors with low initial gas fractions ($f_{\\rm gas}=0.2$) starting below the relations evolve onto the relations, progenitors on the relations stay there and progenitors above the relations evolve towards the relations, but still remaining above them. Progenitors with high initial gas fractions ($f_{\\rm gas}=0.8$) evolve above the relations in all cases, with the initial gas fraction thus being the prime source of scatter in the observed relations (see also \\citealt{2007ApJ...669...45H}). The evolution for all mergers is initially slow with the observed relations typically being rapidly established during a relatively short phase centered at the time of final coalescence of the BHs.  The observations of \\citet{2006ApJ...645..900W,2008ApJ...681..925W} indicate that the BHs at high redshifts were typically overmassive by a factor of $\\sim3$ for a fixed $\\sigma$ and $M_{\\rm bulge}$ compared to the local relation. These BHs could plausibly have been formed during a brief quasar phase (e.g. \\citealp{2006ApJS..163....1H}) triggered during the mergers of very gas-rich galaxies, thus resulting in overmassive BH masses for their given $\\sigma$ and $M_{\\rm bulge}$, as seen in our $f_{\\rm gas}=0.8$ simulation series. However, this being the case it is not obvious how these galaxies would evolve onto the local observed relations until the present day. Another binary merger with massive BHs in place does not bring the galaxies onto the relations (O- and A-series), with high gas fractions mergers moving them even further away from the relations. Potentially, the bulge mass could be increased by dry minor mergers (e.g. \\citealp{2009ApJ...699L.178N}), whereas internal secular processes could be responsible for increasing the velocity dispersion. However, it is not obvious how both the velocity dispersion and the bulge mass could be increased at a fixed BH mass. Finally, another possibility could be that some of the BHs undergoing mergers are ejected in a sling-shot effect \\citep{1974ApJ...190..253S}, thus decreasing the total BH mass.  The fact that the BHs starting below the relations evolve onto the relation, whereas the ones above do not, indicates that our simulations rule out the scenario in which overmassive BHs evolve onto the relations through gas-rich mergers, whereas undermassive BHs can evolve onto the relations in galaxy mergers. Thus, given the numerical limitations inherit in our simulations, our results disfavor the picture in which supermassive BHs develop significantly before their parent bulges.  Finally, our current BH accretion and feedback prescription seems to describe the growth of BH during the merger phase adequately. However, it is not obvious that the present description is also valid during the secularly driven low accretion phase after the merger is completed. Some initial steps to improve the prescription have been taken by developing an entirely new momentum driven feedback prescription \\citep{2009arXiv0909.2872D}. Recent observations \\citep{2007MNRAS.382.1415S,2009ApJ...692L..19S} have indicated that there is a time lag of $\\sim0.5 \\ \\rm Gyr$ between the peak of star formation and the onset of AGN activity. Our current model has difficulties in reproducing this and thus developing models that shed light on this discrepancy might also ultimately help us understanding in how, when and why the BHs in galaxies evolve onto the observed relations."
    },
    "0910/0910.5624_arXiv.txt": {
        "abstract": "{} {By studying the photospheric abundances of 4 RV\\,Tauri stars in the LMC, we test whether the depletion pattern of refractory elements, seen in similar Galactic sources, is also common for extragalactic sources. Since this depletion process probably only occurs through interaction with a stable disc, we investigate the circumstellar environment of these sources. } {A detailed photospheric abundance study was performed using high-resolution UVES optical spectra. To study the circumstellar environment we use photometric data to construct the spectral energy distributions of the stars, and determine the geometry of the circumstellar environment, whereas low-resolution Spitzer-IRS infrared spectra are used to trace its mineralogy.} {Our results show that, also in the LMC, the photospheres of RV\\,Tauri stars are commonly affected by the depletion process, although it can differ significantly in strength from source to source. From our detailed disc modelling and mineralogy study, we find that this process, as in the Galaxy, appears closely related to the presence of a stable Keplerian disc. The newly studied extragalactic objects have similar observational characteristics as Galactic post-AGB binaries surrounded by a dusty disc, and are therefore also believed to be part of a binary system. One source shows a very small infrared excess, atypical for a disc source, but still has evidence for depletion. We speculate this could point to the presence of a very evolved disc, similar to debris discs seen around young stellar objects.} {} ",
        "introduction": "\\label{introduction}  With a distance of $\\sim$\\,50\\,kpc \\citep{feast99}, the Large Magellanic Cloud (LMC) is one of the best galaxies to study stellar evolution. The known distance allows us to calculate the luminosity for stellar sources and it is the ideal laboratory to study physical and chemical processes in an environment with a sub-solar metallicity of $Z\\sim 0.3-0.5$\\,Z$_\\odot$ \\citep{westerlund97}. Moreover, it is close enough to allow detailed studies of individual sources using large-aperture ground based telescopes.  Here we focus on a sample RV\\,Tauri stars in the LMC. RV\\,Tauri stars are pulsating evolved stars with a characteristic light curve showing alternating deep and shallow minima. They are located in the high luminosity end of the Population II Cepheid instability strip \\citep{wallerstein02}. The post-AGB status of RV\\,Tauri stars was a long standing debate, but the detection of circumstellar dust around many objects \\citep{jura86}, and the detection of extragalactic pulsators and their large derived absolute luminosities, were in line with the expected evolutionary tracks of post-AGB stars. The first extragalactic RV\\,Tauri stars in the LMC were discovered by the MACHO experiment \\citep{alcock98}.  \\citet{reyniers07a} and \\citet{reyniers07b} performed a chemical study on two LMC  RV\\,Tauri stars, selected from those reported by \\citet{alcock98}, and found that, like in the Galaxy \\citep{vanwinckel03}, post-AGB stars are chemically much more diverse than previously anticipated. One of the stars, MACHO\\,47.2496.8, proved to be strongly enhanced in s-process elements, in combination with a very high carbon abundance (C/O\\,$>$\\,2 and $^{12}$C/$^{13}$C\\,$\\approx$\\,200). The metallicity of [Fe/H]\\,$=$\\,$-1.4$ is surprisingly low for a field LMC star. The s-process enrichment is large: the light s-process elements of the Sr-peak are enhanced by 1.2 dex compared to iron ([ls/Fe]\\,$=$\\,$+1.2$), while for the heavy s-process (Ba-peak) elements, an even stronger enrichment is measured: [hs/Fe]\\,$=$\\,$+2.1$. Lead was not found to be strongly enhanced. The patterns can only be understood assuming a very low efficiency of the $^{13}$C pocket \\citep{bonacic07} which is created during the dredge-up phenomenon and the associated partial mixing of protons into the intershell. It was the first detailed study of the s-process of a post-AGB star in an external galaxy.  Another object, MACHO\\,82.8405.15, turned out to be chemically altered by the depletion phenomenon \\citep{reyniers07a}. Depletion of refractory elements in the photosphere is a chemical process in which chemical elements with a high dust condensation temperature are systematically underabundant \\citep[e.g.][]{maas05,giridhar05}. The special photospheric abundance patterns are the result of gas-dust separation in the circumstellar environment, followed by re-accretion of only the gas, which is poor in refractory elements. The photospheres become deficient in refractories (as Fe, Ca and the s-process elements), while the non-refractories are not affected. The best abundance tracers of the depletion phenomenon are the Zn/Fe and S/Ti ratios because the elements involved in either ratio have a similar nucleosynthetic formation channel, but have very different condensation temperatures. Intrinsically Fe-poor objects have [Zn/Fe] and [S/Ti] close to solar, which is not the case for depleted objects. With [Fe/H]\\,$=$\\,$-$2.6, in combination with [Zn/Fe]\\,$=$\\,$+$2.3 and [S/Ti]\\,$=$\\,$+$2.5, in MACHO\\,82.8405.15, there is no doubt that the depletion affected the photosphere of this LMC star very strongly. The very low abundances, as well as the clear correlation with condensation temperature, are shown in Figure~\\ref{depletie}, which is a reproduction of the figure in \\citet{reyniers07a}.  Photospheric depletion is surprisingly common in Galactic post-AGB stars \\citep[e.g.][and references therein]{giridhar05,maas05}. In almost all depleted post-AGB objects, there is observational evidence that a stable circumbinary disc is present \\citep{deruyter06, vanwinckel07}. The discs are very compact \\citep[e.g.][]{deroo06, deroo07c} and very likely only found around binary post-AGB stars \\citep[e.g.][and references therein]{vanwinckel09}.  Also in their infrared spectra, the RV\\,Tauri stars show unique spectral features. One of the best studied Galactic RV\\,Tauri stars is AC\\,Her \\citep{molster99}. Infrared studies with ISO have shown very strong crystalline silicate bands from 10--50\\,$\\mu$m. Recent studies with Spitzer show that this strong crystallinity is commonly observed, also in other post-AGB binary sources \\citep{gielen08}.    Thanks to the efficient infrared detectors of the Spitzer satellite, very sensitive infrared observations allow us to probe for circumstellar dust, even around individual objects in external galaxies. The SAGE (Surveying the Agents of a Galaxy's Evolution) Spitzer LMC survey \\citep{meixner06} mapped the LMC, using all photometric bands of the Spitzer IRAC (InfraRed Array Camera, \\citealp{fazio04}) and MIPS (Multiband Imaging Photometer for Spitzer, \\citealp{rieke04}) instruments. This survey resulted in the detection of over 4 million sources. Thanks to the release of this database, we found that the LMC RV\\,Tauri stars as discovered with the MACHO experiment in the visible, indeed have infrared excesses with very similar SED shapes as many Galactic post-AGB binaries (see Fig.~\\ref{kleurkleur}). Since only those LMC sources with strong enough optical fluxes were selected, some observational bias exists to stars where we see the disc more face-on, since edge-on disc would obscure the star too much.    \\begin{figure} \\centering \\resizebox{10cm}{!}{\\includegraphics{12982fg1.ps}} \\caption{Colour-colour diagram indicating the Galactic post-AGB sources with discs (grey circles) and extragalactic RV\\,Tauri stars as presented by \\citet{alcock98} (black circles). The grey dots represent the LMC objects as found by the SAGE-LMC survey \\citep{meixner06}.} \\label{kleurkleur} \\end{figure}    \\begin{table*} \\caption{List of stellar parameters for our sample sources.} \\label{stelpar} \\vspace{0.cm} \\hspace{-.5cm} \\begin{tabular}{lllcrrccccc} \\hline \\hline Name & other name & $\\alpha$ (J2000) & $\\delta$ (J2000)  & $T_{\\rm eff}$ & $\\log g$ & [Fe/H] &  $E(B-V)_{tot}$ & $L_{\\rm IR}/L_*$ & $L_*$     & $P$\\\\ & &(h m s)          & ($^\\circ$ ' '')   & (K)           & (cgs)    &        &                 & (\\%)             & L$_\\odot$ & (days)\\\\ \\hline MACHO\\,79.5501.13 & J051418.1-691235 & 05 14 18.1 & -69 12 34.9 & 5750 & 0.5 & -2.0 & $0.14\\pm0.01$ & $60\\pm3$ & $5000\\pm500$ & 48.5\\\\ & HV\\,915   &&&&&&&&&\\\\ MACHO\\,81.8520.15 & J053254.5-693513 & 05 32 54.5 & -69 35 13.2 & 6250 & 1.0 & -1.5 & $0.27\\pm0.03$ & $0\\pm1$  & $4200\\pm500$ & 42.1\\\\ MACHO\\,81.9728.14 & J054000.5-694214 & 05 40 00.5 & -69 42 14.6 & 5750 & 1.5 & -1.0 & $0.05\\pm0.02$ & $53\\pm3$ & $4200\\pm500$ & 47.1\\\\ MACHO\\,82.8405.15 & J053150.9-691146 & 05 31 51.0 & -69 11 46.4 & 6000 & 0.5 & -2.5 & $0.05\\pm0.01$ & $84\\pm3$ & $4000\\pm500$ & 46.5\\\\ \\hline \\end{tabular} \\begin{footnotesize} \\begin{flushleft} Note: The name, equatorial coordinates $\\alpha$ and $\\delta$ (J2000), effective temperature $T_{\\rm eff}$, surface gravity $\\log g$ and metallicity [Fe/H] of our sample stars. Also given are the total reddening $E(B-V)_{tot}$, the energy ratio $L_{\\rm IR}/L_*$, the calculated luminosity (computed by integrating the dereddened photosphere), assuming a distance of $d=50$\\,kpc, and the period $P$ as given in \\citep{alcock98}. \\end{flushleft} \\end{footnotesize} \\end{table*}  \\begin{table} \\caption{\\label{logUves}Log of the UVES observations and the obtained final S/N at a given spectral band. } \\centering \\vspace{0.5cm} \\hspace{0.cm} \\begin{tabular}{lrrrr} \\hline\\hline\\rule[0mm]{0mm}{3mm} Star               & Exp. Time & Wavelength       & S/N    \\\\ &   (sec.)    &  (nm)              &         \\\\  \\hline MACHO\\,79.5501.13  & 3600      &  375.8 -- 498.3  & 70      \\\\ &           &  670.5 -- 1008.4 & 100      \\\\ MACHO\\,81.8520.15  & 3600      &  478.0 -- 680.8  & 70    \\\\ MACHO\\,81.9728.14  & 3600      &  478.0 -- 680.8  & 65 \\\\ \\hline \\end{tabular} \\begin{footnotesize} \\begin{flushleft} Note: The S/N is measured in the middle of the spectral window covered. \\end{flushleft} \\end{footnotesize} \\end{table}  In this contribution we focus on the MACHO RV\\,Tauri objects in the LMC as detected by the MACHO experiment \\citep{alcock98}, and we connect the chemical studies based on high-resolution optical spectroscopy to the SED energetics and the infrared spectra obtained by Spitzer. Prior to this study, only 1 depleted post-AGB star in the LMC was known. Here we analyse the abundances of 4 more similar sources, testing if depletion is also a common process in the RV\\,Tauri stars of the LMC. As the distance to the LMC is known, we are able to discuss the evolutionary status of these extragalactic sources using their accurate position in the H-R diagram, which sofar has been impossible for Galactic post-AGB stars.  This research was possible thanks to the SAGE-Spec international program (\\textit{http://sage.stsci.edu/}). The goal of this large spectroscopic infrared program is to complement the wealth of photometric data from the SAGE photometric survey, with an extensive spectroscopic follow-up programme using the infrared IRS spectrograph aboard of Spitzer (Kemper et al., 2009, submitted). The main goal of the survey is to determine the composition, origin and evolution, and observational characteristics of interstellar and circumstellar dust and its role in the LMC. A total of about 200 stars, in all stages of stellar evolution, HII and diffuse regions were observed in Spitzer-IRS\tand MIPS-SED mode. ",
        "conclusions": "Clearly the formation of stable dusty discs and the significant feedback of this disc on the central star is not exclusive to our Galaxy alone. Four out of five LMC RV\\,Tauri objects, for which we have high-resolution data, are found to be strongly affected by the depletion process in which the atmospheres became poor in refractory elements. Moreover, the infrared colours and spectral data show that three sources are surrounded by a highly processed, stable disc in which dust processing has been efficient. In one source PAH particles are formed, while there is no other evidence for intrinsic carbon enhancement in the photosphere. This object is intrinsically the most metal poor object of the sample. For MACHO\\,81.8520.15 we only found a small dust excess at 8\\,$\\mu$m and we could interpret this as evidence for strong disc evolution. In this star, even the Zn abundance is affected by depletion which means that the gas-dust separation occurred at low temperatures.  Overall the RV\\,Tauri stars in the LMC display many characteristics of their Galactic peers. Whether these extragalactic stars also reside in a binary system remains unclear with the available data at hand. The low magnitude of these stars does not allow for a long-term radial velocity monitoring programme but, given the strong similarities to the Galactic binaries, a binary evolutionary channel is very likely needed to understand the RV\\,Tauri stars in the LMC as well."
    },
    "0910/0910.2468_arXiv.txt": {
        "abstract": "We present the results of an X-ray mass analysis of the early-type galaxy NGC~4636, using \\Chandra\\ data. We have compared the X-ray mass density profile with that derived from a dynamical analysis of the system's globular clusters (GCs). Given the observed interaction between the central active galactic nucleus and the X-ray emitting gas in NGC~4636, we would expect to see a discrepancy in the masses recovered by the two methods. Such a discrepancy exists within the central $\\sim$10\\,kpc, which we interpret as the result of non-thermal pressure support or a local inflow. However, over the radial range $\\sim$\\,10--30\\,kpc, the mass profiles agree within the 1\\,$\\sigma$ errors, indicating that even in this highly disturbed system, agreement can be sought at an acceptable level of significance over intermediate radii, with both methods also indicating the need for a dark matter halo. However, at radii larger than 30\\,kpc, the X-ray mass exceeds the dynamical mass, by a factor of 4--5 at the largest disagreement. A Fully Bayesian Significance Test finds no statistical reason to reject our assumption of velocity isotropy, and an analysis of X-ray mass profiles in different directions from the galaxy centre suggests that local disturbances at large radius are not the cause of the discrepancy. We instead attribute the discrepancy to the paucity of GC kinematics at large radius, coupled with not knowing the overall state of the gas at the radius where we are reaching the group regime ($>$30\\,\\kpc), or a combination of the two. ",
        "introduction": "\\label{sec:intro} The current paradigm of galaxy formation describes how galaxies form embedded in massive dark matter halos. Whereas the measurement of rotation curves can be successfully applied to late-type galaxies to infer the presence of this dark matter \\citep[see e.g.][for a review]{sofue01}, this cannot be employed in early-type galaxies as their stars and gas are not supported by rotation. Therefore, different methods must be invoked to measure the galaxy mass.  It has long been known that early-type galaxies contain hot ($\\sim$10$^{6}$K) X-ray emitting gas \\citep{forman85}, the temperature and density of which allow the determination of the total gravitating mass, assuming hydrostatic equilibrium and spherical symmetry. This approach has proved successful at yielding meaningful mass profiles for early-type galaxies \\citep[e.g.][]{osullivan04b,fukazawa06,humphrey06,zhang07}. The effect of the assumption of spherical symmetry has been addressed in the case of galaxy clusters, indicating that although compression and elongation along the line-of-sight can under or over-estimate the central mass respectively, this is only a small effect at large radius \\citep{piffaretti03}. However, the validity of the intrinsic assumption of hydrostatic equilibrium has been questioned with specific reference to early-type galaxies \\citep{diehl06}. NGC~4636 presents an ideal test-bed in this respect, as it is a highly disturbed system, with evidence of bubbles and shocks caused by previous AGN outbursts \\citep{jones02,ohto03,osullivan05b}.  The use of dynamical tracers of the gravitational potential is also a well-established method to recover the kinematics of bound systems, both on the scale of globular clusters \\citep[e.g.][in M15]{gebhardt00}, and for the Galaxy itself \\citep{chakrabarty01,genzel00,ghez98}. The use of globular clusters (GCs) as tracers of the potential has been particularly successful in recovering mass profiles of nearby elliptical galaxies \\citep[examples include][]{romanowsky01,cote03,bergond06,schuberth06,woodley07}. Similarly, dedicated surveys of planetary nebulae in early-type galaxies can also be used to derive the distribution of matter \\citep{douglas07}, although care is required to avoid complications from distinct populations of planetary nebulae, which have been seen for example in the galaxy NGC 4697 \\citep{sambhus06}. These approaches involve solving the Jeans equations under the assumption of spherical symmetry to determine the galaxy mass. Interestingly, a study of three early-type galaxies using planetary nebulae kinematics, \\citet{romanowsky03} concluded a significant lack of dark matter in these systems. However, \\citet{dekel05} showed these data to be consistent with a massive dark halo when more radial orbits were considered. This highlights the mass-anisotropy degeneracy present in this approach, which can be broken by considering higher order velocity moments \\citep{lokas07}. A further systematic effect is the assumption of spherical symmetry. In the case where the galaxy is flattened along the line-of-sight, its mass can be under-estimated if the system is assumed to be spherically symmetric \\citep{magorrian01}.  As both X-ray and dynamical methods have their own intrinsic assumptions, the most robust constraints can be placed on the mass profiles of early-type galaxies when different methods are compared. Indeed, there is currently an emerging attempt to use different techniques in a complementary manner \\citep[][]{romanowsky08,churazov08,samurovic06,bridges06}; additionally, this approach improves our understanding of the systematics involved in each method. Recent work by \\citet{churazov08} explored in detail the comparison between X-ray and optically derived profiles for M87 and NGC 1399, finding agreement between the methods at the 10--20\\,\\% level when looking at the gravitational potential. However, both of these systems reside at the centres of clusters, M87 being the centre of Virgo, and NGC 1399 the centre of Fornax, and in both cases, the measurement of the potential is probing the cluster potential. In so-called `normal' elliptical galaxies, the situation is much less certain. For example, in the galaxy NGC 3379, \\citet{pellegrini06} require an outflow of the X-ray emitting gas to bring the X-ray results into agreement with the optically derived results. Only by the study of more systems with multiple approaches will we be able to reconcile the observed discrepancies. This is successful on a local scale, as individual GCs and/or planetary nebulae need to be resolved, limiting the distance to which these observations can be made. Investigating the wider properties of the dark matter halos of elliptical galaxies will require techniques such as stacked lensing \\citep[e.g.][]{sheldon04,hoekstra05,kleinheinrich06,koopmans06,mandelbaum06,ferreras08}.  The layout of the paper is as follows. Section \\ref{sec:N4636} describes the basic properties of NGC~4636 and Section \\ref{sec:x-ray} describes our method for extracting high resolution mass profiles from \\Chandra\\ X-ray data. In Section \\ref{sec:results} we present our results and comparison to the GC analysis of \\citet{chakrabarty08}, and in Section \\ref{sec:discuss} we discuss the implications of our results. ",
        "conclusions": "We present an X-ray mass analysis of the early-type galaxy NGC~4636 using \\Chandra\\ data, under the assumptions of spherical symmetry and hydrostatic equilibrium. The integrity of the latter assumption has been questioned with reference to early-type galaxies \\citep{diehl06}, and it is because of the observed disturbances in the gas in NGC~4636 \\citep[e.g.][]{jones02,ohto03,osullivan05b} that we chose to study this object, in an effort to assess the impact of this assumption on the recovered mass profile. We find that the treatment of the abundance gradient in the X-ray analysis can significantly affect the recovered mass profile at all radii.  We have compared the X-ray mass density profile with that recovered from a dynamical analysis of the system's globular clusters (GCs), presented by \\citet{chakrabarty08}. Inside 10\\,\\kpc, the dynamical mass estimate exceeds the X-ray mass estimate. The gas in this region is highly disturbed, and we postulate the cause of the disagreement to be a localised inflow of gas, or a contribution of non-thermal pressure support.  The mass density profiles over the range $\\sim$10--30\\,\\kpc\\ are consistent within 1\\,$\\sigma$, indicating that even in this highly disturbed system, the recovered X-ray mass is consonant with that recovered from an independent method over intermediate radii. However, outside 30\\,\\kpc, the X-ray mass estimate exceeds the dynamical mass estimate, by a factor of 4--5 times at its greatest disagreement. Examining the anisotropy of the GCs, we find no statistical reason to reject our assumption of isotropy. The GC analysis is model-independent, so is not limited by the method, but the paucity of measured GC kinematics outside 7.5$^{\\prime}$ means that the success of this method at large radius is limited by the data.  We test the assumption of hydrostatic equilibrium in our X-ray analysis, finding that local disturbances at large radii do not account for the observed discrepancy. At this radius, the group gas contribution is important in this system \\citep{osullivan05b}, and the overall state of the gas at this radius is uncertain. Mapping the X-ray properties to a larger radius using \\XMM\\ would help to model the group emission, but is beyond the scope of the current paper.  The X-ray and dynamical mass analysis methods both indicate the need for a dark matter halo in this system, and provide a useful comparison within 30\\,\\kpc. It is through the comparison of independent approaches that the most robust constraints will be placed on the mass distribution of early-type galaxies, but we conclude that the limiting factors in such a comparison to large radius (outside 30\\,\\kpc) are data quality in the case of the GC kinematics, knowledge of the overall state of the gas as we reach the group regime in the case of the X-ray analysis, or a combination of the two."
    },
    "0910/0910.0976_arXiv.txt": {
        "abstract": "Long-term trends in the solar spectral irradiance are important to determine the impact on Earth's climate. These long-term changes are thought to be caused mainly by changes in the surface area covered by small-scale magnetic elements. The direct measurement of the  contrast to determine the impact of these small-scale magnetic elements is, however, limited to a few wavelengths, and is, even for space instruments,  affected by scattered light and instrument defocus. In this work we calculate emergent intensities from 3-D simulations of solar magneto-convection and validate the outcome by comparing with observations from Hinode/SOT. In this manner we aim  to construct the contrast at wavelengths ranging from the NUV to the FIR. ",
        "introduction": "Variations in the long-term spectral solar irradiance, i.e. the changes in the Sun's brightness at a certain wavelength, are significant for their effects on Earth's atmospheric temperature and composition. Magnetic fields concentrated in small structures  influence these long-term changes, with increasing relevance towards the solar limb. Here, we compare observed and simulated  intensity contrasts and investigate their behaviour from disk centre to the limb in 3D MHD simulations with different average vertical magnetic fields. ",
        "conclusions": "This analysis will allow us to determine the network and facular contrast as a function of magnetic flux, limb angle, and wavelength. We will then be able to remove the single free parameter in the SATIRE model (e.g. Krivova et al. 2003) and help improve irradiance reconstructions."
    },
    "0910/0910.5886_arXiv.txt": {
        "abstract": "For axisymmetric models for coronal loops the relationship between the bifurcation points of magnetohydrodynamic (MHD) equilibrium sequences and the points of linear ideal MHD instability is investigated imposing line-tied boundary conditions. Using a well-studied example based on the Gold-Hoyle equilibrium, it is demonstrated that if the equilibrium sequence is calculated using the Grad-Shafranov equation, the instability corresponds to the second bifurcation point and not the first bifurcation point because the equilibrium boundary conditions allow for modes which are excluded from the linear ideal stability analysis. This is shown by calculating the bifurcating equilibrium branches and comparing the spatial structure of the solutions close to the bifurcation point with the spatial structure of the unstable mode. If the equilibrium sequence is calculated using Euler potentials the first bifurcation point of the Grad-Shafranov case is not found, and the first bifurcation point of the Euler potential description coincides with the ideal instability threshold. An explanation of this results in terms of linear bifurcation theory is given and the implications for the use of MHD equilibrium bifurcations to explain eruptive phenomena is briefly discussed. ",
        "introduction": "Magnetohydrodynamic instabilities of coronal loops are since a long time discussed as one of the main theoretical explanations for solar flares, in particular compact loop flares ({\\it e.g.} \\opencite{priest82}). Traditionally, investigations of MHD instabilities of coronal loops model these loops as straight cylindrical flux tubes of finite length with line tied boundary conditions at the `photospheric' ends of the flux tubes ({\\it e.g.} \\opencite{raadu72}; \\opencite{hood:priest79}, \\citeyear{hood:priest81}; \\opencite{einaudi:vanhoven83}; \\opencite{velli:etal90}). Such a set-up allows for a wide variety of relatively simple equilibrium configurations, hence explaining its popularity.  The stability of equilibrium configurations of the above mentioned type has been studied for several decades using the methods of linear MHD stability analysis ({\\it e.g.} \\opencite{raadu72}; \\opencite{hood:priest79}, \\citeyear{hood:priest81}; \\opencite{einaudi:vanhoven83}; \\opencite{velli:etal90}; \\opencite{debruyne:hood89}, \\citeyear{debruyne:hood92}; \\opencite{mikic:etal90}; \\opencite{hood:etal94}; \\opencite{vanderlinden:hood98}, \\citeyear{vanderlinden:hood99}). In recent years the investigations have been extended into the nonlinear regime using large-scale MHD simulations ({\\it e.g.} \\opencite{longbottom:etal96}; \\opencite{baty:heyvaerts96}; \\opencite{baty97a}, \\citeyear{baty97b}, \\citeyear{baty00a}, \\citeyear{baty00b}; \\opencite{lionello:etal98}; \\opencite{arber:etal99}; \\opencite{gerrard:etal01}; \\opencite{browning:vanderlinden03}; \\opencite{browning:etal08}; \\opencite{hood:etal09}).  In the present contribution we want to investigate the stability of line-tied coronal loop models from a different point of view. The flux tube equilibria used to model coronal loops all depend on one or more parameters representing quantities like the magnetic twist or the plasma beta. Many investigations study how the linear stability of the loops changes as one (or more) of these equilibrium parameters vary.  The systematic variation of one or several parameters of an equilibrium defines an equilibrium sequence, and a point of linear instability should correspond to a bifurcation point of the equilibrium sequence and vice versa. It has to be kept in mind, however, that magnetostatic equilibria are usually calculated by solving a mathematically reduced set of equations. It is not at all clear whether there is really a one-to-one correspondence between points of linear instability and bifurcation points, in particular if line-tied boundary conditions are imposed as in models of coronal loops.  In the present paper we shall investigate the question whether the points of linear instability of rotationally symmetric straight line-tied flux tubes have a one-to-one correspondence with the bifurcation points of equilibrium sequences.  We shall use two different ways of calculating the equilibrium sequences, namely Grad-Shafranov theory and Euler potentials, and we shall, for simplicity, investigate only axisymmetric instabilities and bifurcations. A particularly well-studied equilibrium class \\cite{gold:hoyle60} will be used to carry out this investigation, mainly because results of linear stability investigations for this equilibrium class are readiliy available in the literature ({\\it e.g.} \\opencite{hood:priest79}, \\citeyear{hood:priest81}; \\opencite{mikic:etal90}; \\opencite{debruyne:hood92}).  In Section \\ref{basic} the basic equilibrium theory and those parts of the theory of linear MHD stability needed in this paper are discussed. The following Section \\ref{numerics} presents a brief outline of the numerical method used to calculate the equilibrium sequences and to determine their bifurcation points and bifurcating branches. The results of these calculations are given in Section \\ref{results} and discussed in Section \\ref{discussion}. The paper closes with a summary in Section \\ref{summary}. ",
        "conclusions": "\\label{summary}  We have investigated the relationship between MHD bifurcation and linear stability for a class of axisymmetric straight flux tubes under line-tying boundary conditions. For simplicity we only considered rotationally symmetric perturbations, allowing only for sausage modes. We have used two different ways of calculating the equilibrium sequences including bifurcating branches - one approach uses the Grad-Shafranov equation, the other approach uses Euler potentials. It turns out that only the Euler potential case shows a one-to-one correspondence between the first bifurcation point and the linear instability threshold for the sausage mode. The Grad-Shafranov case shows an additional bifurcation which does not correspond to the instability threshold under line-tying boundary conditions. This difference can be explained by the different constraints imposed on the bifurcating equilibrium branches in the Grad-Shafranov and the Euler potential cases.  Furthermore, even though the second bifurcation point of the Grad-Shafranov case coincides with the first bifurcation point of the Euler potential case and the linear instability threshold, the structure of the bifurcation diagrams differ considerably between the Grad-Shafranov and the Euler potential case. The reason for this is not yet clear, but is probably also due to the difference in boundary conditions. In any case this difference has implications for the stability of the bifurcating equilibrium branches (see {\\it e.g.} \\opencite{iooss:joseph80}) and is therefore important to decide whether the system is able to find a new equilibrium (in the present case a new axisymmetric equilibrium) if one would consider an imaginary process driving the flux tube across the instability threshold.  The present investigation is a preparation for studying equilibrium sequences of magnetic  flux tubes and other solar magnetic structures together with their bifurcations in three dimensions. Preliminary steps have already been made (see {\\it e.g.} \\opencite{romeou:neukirch02}) and more detailed investigations are planned for the future.   \\begin{acks}  The authors thank Alan Hood for useful discussions. T. Neukirch acknowledges support by STFC and by the European Commission through the SOLAIRE Network (MTRN-CT-2006-035484). Z. Romeou gratefully acknowledges financial support provided through the European Community's Training and Mobility of Researchers Programme by a Marie-Curie Fellowship and through the European Community's Human Potential Programme under contract HPRN-CT-2000-00153, PLATON. The authors also acknowledge partial support by the British Council ARC Programme.  \\end{acks}  \\begin{appendix}   Whereas the first bifurcating branch in the Grad-Shafranov case exists for values of $\\lambda$ which are both bigger and smaller than the $\\lambda$ at the bifurcation point, the second bifurcating branch exists only for $\\lambda$ smaller than the bifurcation $\\lambda$. This fact can be explained by using standard bifurcation theory to calculate the structure of the bifurcating branches close to the bifurcation points. The argument is actually independent of the form of the functions $p(A)$ and $b_\\phi(A)$. The qualitative structure of the bifurcation diagram will thus be the same even if $p(A)$ and $b_\\phi(A)$ are changed as long as the fundamental branch consists of solutions which depend only on the radial coordinate $r$.  We start by writing the Grad-Shafranov equation in the form \\begin{equation} G(A,\\lambda,r)=-r\\nabla\\cdot\\left(\\frac{1}{r^2}\\nabla A\\right) - N(A,\\lambda,r) = 0, \\label{gsgeneral} \\end{equation} where the function $N(A,r,\\lambda)$ summarizes the nonlinear part of the Grad-Shafra\\-nov equation given by $p(A,\\lambda)$ and $b_\\phi(A,\\lambda)$. For the present paper $p$ and $b_\\phi$ are given by Equations (\\ref{ghpa}) and (\\ref{ghbpa}). For the following argument, however, the exact form of $N(A,r,\\lambda)$ is irrelevant, as long as it is analytic in $A$ and $\\lambda$ at the bifurcation points we want to investigate. We will not give here any details of the mathematical background which can be found for example in \\inlinecite{hesse:ks86} and\\inlinecite{hesse:kiessling87}. These papers treat slighly different bifurcation problems, but we will be using the same technique.  Let $\\lambda^*$ be the value of $\\lambda$ at either of the bifurcation points and let $A_0=A_0(\\lambda^*)$ be the solution of Equation (\\ref{gsgeneral}) at the bifurcation point. To calculate the bifurcating branch we expand $\\lambda$ and $A$ as \\begin{eqnarray} \\lambda &=& \\sum_{k=0}^\\infty \\epsilon^k \\lambda_k ,\\label{lambdaexp} \\\\ A &=& \\sum_{k=0}^\\infty \\epsilon^k A_k        ,      \\label{Aexp} \\end{eqnarray} where $\\lambda_0=\\lambda^*$ and $A_0$ as above. Since $G$ is analytic in both $A$ and $\\lambda$ for $\\lambda > 0$ we can expand Equation (\\ref{gsgeneral}) in a power series in $\\epsilon$: \\begin{equation} 0 = \\sum_{k=0}^\\infty \\frac{1}{k !}\\frac{d^k }{d\\epsilon^k} G(A(\\epsilon),\\lambda(\\epsilon),r )\\Big|_{\\epsilon=0} \\epsilon^k . \\end{equation} As each power of $\\epsilon$ must satisfy this equation independently we obtain \\begin{equation} \\frac{d^k }{d\\epsilon^k} G(A(\\epsilon),\\lambda(\\epsilon),r )\\Big|_{\\epsilon=0} = 0, \\qquad k=0,1,2,\\ldots\\; . \\label{Gexp} \\end{equation} Obviously, the lowest order equation \\begin{equation} G(A_0(\\lambda^*),\\lambda^*,r) = 0 \\end{equation} is just the Grad-Shafranov equation at the bifurcation point and therefore trivially satisfied.  For the discussion of the higher order equations we first have to look at the boundary conditions the $A_k$ have to satisfy. The boundary condtion $A_b$ for $A(r,z,\\lambda)$ is given by the fundamental branch solution $A_b(r,z,\\lambda)=A_0(r,z,\\lambda)$ (the Gold-Hoyle solution in the present paper). Therefore we can extend $A_b$ into the domain. Using the expansion (\\ref{lambdaexp}) in $A_b(r,z,\\lambda(\\epsilon))$ we can see that the boundary condition each of the $A_k$ in Equation (\\ref{Aexp}) has to satisfy is given by \\begin{equation} A_b^{(k)} = \\frac{1}{k !}\\frac{d^k}{d \\epsilon^k} A_0(r,z,\\lambda(\\epsilon)) \\Big|_{\\epsilon=0}. \\end{equation} Note that $A_b^{(k)}$ satifies the same Equation (\\ref{Gexp}) as $A_k$.  For $O(\\epsilon)$ we get from Equation (\\ref{Gexp}) \\begin{equation} G_A(A_0(\\lambda^*),\\lambda^*, r) A_1 + G_\\lambda(A_0(\\lambda^*),\\lambda^*, r) \\lambda_1 =0, \\end{equation} with \\begin{eqnarray} G_A A_1 &=& -r\\nabla\\cdot\\left(\\frac{1}{r^2}\\nabla A_1\\right) - \\frac{\\partial N}{\\partial A}(A_0,\\lambda^*,r) A_1 ,\\\\ G_\\lambda &=&\\frac{\\partial N}{\\partial \\lambda}(A_0,\\lambda^*,r). \\end{eqnarray} As mentioned above $A_b^{(1)}$ satisfies the same equation as $A_1$ and therefore the function \\begin{equation} A_1^\\prime = A_1 -A_b^{(1)} \\end{equation} satisfies $G_A A_1^\\prime =0$ or explicitely \\begin{equation} -r\\nabla\\cdot\\left(\\frac{1}{r^2}\\nabla A_1^\\prime\\right) - \\frac{\\partial N}{\\partial A}(A_0,\\lambda^*,r) A_1^\\prime =0 \\label{firstorder} \\end{equation} with $A_1^\\prime = 0$ on the boundaries. Since all coefficients of Equation (\\ref{firstorder}) depend only on $r$ its solution can be obtained by separation of variables with the general form of $A_1^\\prime$ being \\begin{equation} A_1^\\prime(r,z) = F_n(r) \\sin(n\\pi z/L), \\qquad n= 1,2,3,\\ldots \\;. \\label{a1modes} \\end{equation} The exact form of $F_n(r)$ is of no importance for the following argument.  If we want to calculate $\\lambda_1$, we have to go to the next order ($O(\\epsilon^2)$) of the expansion, giving \\begin{eqnarray} & &-r\\nabla\\cdot\\left(\\frac{1}{r^2}\\nabla A_2\\right)-\\frac{\\partial N}{\\partial A} A_2 -\\frac{1}{2}\\frac{\\partial^2 N}{\\partial A^2} A_1^2 - \\frac{\\partial^2 N}{\\partial A \\partial \\lambda} A_1 \\lambda_1 \\nonumber \\\\ & &\\mbox{ \\hspace{5cm}}-\\frac{1}{2}\\frac{\\partial^2 N}{\\partial \\lambda^2} \\lambda_1^2 - \\frac{\\partial N}{\\partial \\lambda} \\lambda_2 =0 \\end{eqnarray} where all derivatives of $N(A,\\lambda,r)$ are evaluated at the bifurcation point ($\\epsilon=0$). Similarly to $A_b^{(1)}$ at $O(\\epsilon)$, $A_b^{(2)}$ satisfies the same equation as $A_2$. We define \\begin{equation} A_2^\\prime=A_2-A_b^{(2)} \\end{equation} which obeys the equation \\begin{equation} G_A A_2^\\prime = \\frac{1}{2}\\frac{\\partial^2 N}{\\partial A^2} ({A_1^\\prime}^2 + 2 A_b^{(1)} A_1^\\prime) + \\lambda_1\\frac{\\partial^2 N}{\\partial A \\partial \\lambda} A_1^\\prime . \\label{secondorder} \\end{equation} By Fredholm's alternative the right hand side of Equation (\\ref{secondorder}) has to be orthogonal to $A_1^\\prime$, {\\it i.e.} \\begin{equation} \\int_0^L\\int_0^{r_{max}}\\left( \\frac{1}{2}\\frac{\\partial^2 N}{\\partial A^2} ({A_1^\\prime}^2 + 2 A_b^{(1)} A_1^\\prime) + \\lambda_1\\frac{\\partial^2 N}{\\partial A \\partial \\lambda} A_1^\\prime \\right) A_1^\\prime rdrdz = 0 . \\label{fredha} \\end{equation} To proceed we assume in agreement with the Gold-Hoyle solution that the function $A_b^{(1)}$ has the form \\begin{equation} A_b^{(1)} = \\lambda_1 f_b(r) \\end{equation} where $f_b(r)$ is left unspecified here. Equation (\\ref{fredha}) can then be used to calculate $\\lambda_1$ in the form \\begin{eqnarray} \\lefteqn{ \\lambda_1 \\int_0^L\\int_0^{r_{max}}\\left( \\frac{\\partial^2 N}{\\partial A^2} f_b(r) + \\frac{\\partial^2 N}{\\partial A \\partial \\lambda}\\right) {A_1^\\prime}^2 r dr dz  =   }   \\nonumber \\\\ & & \\mbox{\\hspace{4cm}}-\\int_0^L\\int_0^{r_{max}} \\frac{1}{2}\\frac{\\partial^2 N}{\\partial A^2} {A_1^\\prime}^3 r dr dz . \\label{l1equat} \\end{eqnarray} The double integral on the right hand side of Equation (\\ref{l1equat}) can be split into two separate integrations over $r$ and $z$, since the integrand depends on $z$ only through $A_1^\\prime$. As $A_1^\\prime$ has the form (\\ref{a1modes}), the integral over $z$ is given by \\begin{equation} \\int_0^L \\sin^3(n\\pi z/L) dz = -\\frac{L}{n\\pi}(\\cos n\\pi -1) +\\frac{L}{3n\\pi}(\\cos^3 n\\pi -1) , \\end{equation} which vanishes for all even $n$. As the integral on the left hand side of Equation (\\ref{l1equat}) is nonzero, this implies that for even $n$ (and in particular for $n=2$) $\\lambda_1$ vanishes. The bifurcation at bifurcation points with modes having even $n$ is therefore quadratic.  This explains the structure of the bifurcation diagram in Figure \\ref{bifdiaggs7}, because the first bifurcation obviously corresponds to the $A_1^\\prime$ for $n=1$, whereas the second bifurcation corresponds to $n=2$. Therefore the structure of the first branch close to the bifurcation point is given by \\begin{eqnarray} \\lambda & = &\\lambda^* + \\epsilon \\lambda_1 + \\ldots, \\label{lambdabif1} \\\\ A       & = & A_0 (r,z,\\lambda^*) + \\epsilon A_1(r,z,\\lambda^*) + \\ldots.\\label{Abif1} \\end{eqnarray} The slope of the bifurcating branch at the bifurcation point is determined by $\\lambda_1 \\ne 0$ in this case and it is obvious that the bifurcating branch exists for both $\\lambda > \\lambda^*$ ($\\epsilon \\lambda_1 > 0$) and $\\lambda < \\lambda^*$ ($\\epsilon \\lambda_1 < 0$).  Close to the second bifurcation point we have \\begin{eqnarray} \\lambda & = &\\lambda^* + \\frac{1}{2}\\epsilon^2 \\lambda_2 + \\ldots, \\label{lambdabif2}\\\\ A       & = & A_0 (r,z,\\lambda^*) + \\epsilon A_1(r,z,\\lambda^*) + \\ldots,  \\label{Abif2} \\end{eqnarray} because here $\\lambda_1$ vanishes. As the correction to $\\lambda^*$ depends quadratically on $\\epsilon$, positive and negative $\\epsilon$  give the same value of $\\lambda$. This implies that the second bifurcating branch actually consists of {\\em two} branches, one for positive and one for negative $\\epsilon$. Since a change of sign of $\\epsilon$ in Equation (\\ref{Abif2}) corresponds to a simple mirroring of the $\\sin(2\\pi z/L)$ function at the point $z=L/2$, the two branches have exactly the same energies. We remark that since $\\lambda_1=0$ in this case $A_1 = A_1^\\prime$ as the boundary contribution to $A_1^\\prime$ vanishes. The numerical calculations corroborate these results as the same second bifurcation branch is found by the code starting both with negative and positive $\\epsilon$. The only difference between the calculations is the mirroring of the $z$-dependence of $A$ along the bifurcating branch.  \\end{appendix}"
    },
    "0910/0910.5248_arXiv.txt": {
        "abstract": "Microgauss magnetic fields are observed in all galaxies at low and high redshifts. The origin of these intense magnetic fields is a challenging question in astrophysics.  We show here that the natural plasma fluctuations in the primordial universe (assumed to be random), predicted by the Fluctuation-Dissipation-Theorem, predicts $\\sim 0.034~ \\mu G$ fields over $\\sim 0.3$ kpc regions in galaxies. If  the dipole magnetic fields predicted by the Fluctuation-Dissipation-Theorem are  not completely random,  microgauss fields over regions $\\gtrsim 0.34$ kpc are easily obtained. The model is thus a strong candidate for resolving  the problem of the origin of magnetic fields  in $\\lesssim 10^{9}$ years  in high redshift galaxies. ",
        "introduction": "The origin of large-scale cosmic magnetic fields in galaxies and protogalaxies remains   a challenging problem  in astrophysics \\citep{zwe97, kul08, raf08, wid02}. There have been many attempts to explain  the origin of cosmic magnetic  fields. One of the  first popular  astrophysical theories to create seed fields was   the Biermann mechanism \\citep{bie50}. It has been suggested that this mechanism acts in diverse astrophysical systems, such as  large scale structure formation \\citep{pee67,ree72,was78}, cosmological ionizing fronts \\citep{gne00},   star formation and  supernova explosions  \\citep{mir98, han05}.  \\citet{ryu08} made simulations showing that cosmological shocks can create   average magnetic fields of a few $\\mu G$ inside cluster/groups,  $ \\sim 0.1 ~\\mu G$ around clusters/groups, and $\\sim 10~nG$ in filaments.  \\citet{med06} showed that magnetic fields can be produced by collisionless shocks in galaxy clusters and in the intercluster medium (ICM) during large scale structure formation. \\citet{ars09} studied the evolution of magnetic fields  in galaxies coupled with  hierarchical structure formation.   \\citet{ich06}   investigated  second-order couplings between photons and electrons  as a possible origin of  magnetic fields on cosmological scales before the epoch of recombination.  The creation of early magnetic fields generated by cosmological perturbations have also been investigated  \\citep{tak05, tak06, cla01, mae09}.    In our galaxy, the magnetic field is coherent over kpc scales with alternating directions in the arm and inter-arm regions (e.g.,\\citet{kro94, han08}). Such alternations are expected for magnetic fields of primordial origin \\citep{gra01}.  Various observations put upper limits on  the intensity of  a homogeneous primordial magnetic field. Observations of the small-scale cosmic microwave background (CMB) anisotropy yield an upper comoving limit of $4.7 ~nG$ for a homogeneous primordial field \\citep{yam06}. Reionization of the Universe puts upper limits of $0.7-3~ nG$ for a homogeneous primordial field, depending on the assumptions of the stellar population that is responsible for reionizing the Universe \\citep{sch08}. Another upper limit for a homogenous primordial magnetic field is the magnetic Jeans mass $\\sim 10^{10} M_{\\odot} ~(B/3nG)^{3}$ \\citep{sub98, set05}. Thus, if we are investigating the collapse of a $\\sim 10^{7} M_{\\odot}$ protogalaxy, the homogeneous primordial magnetic field must be $< 0.3~ nG$ in order for collapse to occur.  Galactic magnetic fields have been  suggested to  have  evolved in three main stages.  In the first stage,   seed fields were embedded in the protogalaxy.   They may have had   a   primordial origin,  as suggested in this paper. Another possibility is that the seed fields could have been injected  into the protogalaxies by AGN jets, radio lobes, supernovas,  or a combination of the above.  Still another possibility is that  the seed fields may  have been   created by the Biermann battery  during the formation of the protogalaxy.  In the second   stage, the seed  fields were amplified by compression,   shearing flows, turbulent flows, magneto-rotational instabilities,  dynamos or by a combination of the above. In the last stage magnetic fields were ordered by a large scale dynamo \\citep{beck06}.          \\citet{ryu08} investigated the amplification of magnetic fields due to turbulent vorticity created at cosmological shocks during the formation of large scale structures.  A given vorticity $\\omega$ can be characterized by a characteristic velocity $V_{c}$ over a characteristic distance $L_{c}$. \\citeauthor{ryu08} found  that $\\omega$ typically is \\begin{equation} \\omega \\sim 1-3 \\times 10^{-16} s^{-1}, \\end{equation} which corresponds to 10-30 turnovers in the age of the universe. They investigated $L_{c} > $ 1 Mpc $h^{-1}$. We investigate $L_{c} \\simeq$ 200 kpc $h^{-1}$ in  protogalaxies for  a  similar vorticity.   We show that a seed field $0.003 ~nG$ over a comoving $2~ kpc$ region at  \\emph{z} $\\sim 10$,  predicted by the Fluctuation-Dissipation Theorem \\citep{raf08},  amplified by the small scale  dynamo is a good candidate for the origin of magnetic fields in galaxies. \\citet{sub97, sub99} and \\citet{bra00} derived the non-linear evolution equations for the magnetic correlations.   We use their formulation for the small scale  dynamo  and solve the nonlinear  equations numerically.  In \\S II,  we review  the creation of  magnetic fields due to electromagnetic fluctuations in hot dense equilibrium primordial plasmas,  as described in our previous work \\citep{raf08}.  In \\S III, we discuss the   small scale  dynamo and in \\S IV,   the important parameters of the plasma to be used in the calculations.  In \\S V,   we present our results and in \\S VI our conclusions. ",
        "conclusions": "It was shown previously that the magnetic fields,   created immediately after the quark-hadron transition, produce  relatively intense magnetic dipole fields on small scales at $\\emph{z} \\sim 10$ \\citep{raf08}. We show here that   the predicted seed fields of size $\\sim 2$ kpc and intensity $0.003~ nG$ at \\emph{z} $\\sim 10$    can be amplified by a small scale  dynamo  in protogalaxies to intensities  close to observed values.  In the small scale  dynamo studied, we use the  turbulent spectrum given by   \\citet{sub99}.  The characteristic velocity $V_{c}$ and length $L_{c}$,  used in the expression for the  vorticity $V_{c}/L_{c}$,  are  $V_{c} \\simeq 10^{7} cm/s$ and $L_{c} \\simeq$ 200 kpc. This vorticity is comparable to  that found by \\citet{ryu08},   studying   the formation of large scale structures. The length $L_{c} \\simeq 200$ kpc used is a characteristic size of a protogalactic cloud.   The turbulent spectrum used simulates Kolmogorov turbulence \\citep{vai82}.  From  our Figs. 1 and 2,  we find that $M_{L}(\\sim B^{2})$ increases from $\\sim 10^{-23} ~G^{2}$ (corresponding to a magnetic field $B\\sim 3\\times 10^{-12} ~G$  over a region $L\\sim $ 2 kpc) to $M_{L} \\sim 10^{18}~G^{2}$ (corresponding to a field $\\sim 10^{-9}~G$ over a region $L\\sim 2$ kpc)  in $10^{9}$ years. This corresponds to  a $\\sim 6$ e-fold amplification of $B$ in a relatively short time. Collapsing to form galaxies at redshift $\\emph{z}\\sim10$, the density increases by a factor of  $\\sim 200$ and the   magnetic fields  are amplified by a factor of $\\sim 34$. This predicts $0.03~ \\mu G$ fields over 0.34 kpc regions  in galaxies.  If the dipole magnetic fields predicted by the Fluctuation-Dissipation Theorem are  not completely random, microgauss fields over regions $> 0.34$ kpc are easily obtained. The model studied is thus a strong candidate  to explain the $\\mu G$ fields observed in high redshift galaxies.   \\begin{figure} \\centering \\includegraphics[scale=0.55]{f1.eps} \\caption{Values of $M_{L}(G^{2}$) as a function of t (years) and r. Solid black line has the  reference values: $M_{L}(r,0)= 10^{-11} (0.1pc/r)^{3}~G^{2}$,  $L_{c} = 200 ~kpc$, r = 3 kpc,   and $V_{c}=10^{7}~ cm/s$ in Eqs. (36)-(38).   Dashed red line is for    $r= 4 ~kpc$, and the dotted blue line  for  $r = ~5 kpc$.} \\label{f1} \\end{figure} \\begin{figure} \\centering \\includegraphics[scale=0.55]{f2.eps} \\caption{Values of $M_{L}(G^{2}$) as a function of t (years),  varying $V_{c}$. Solid black line has the  reference values in  Fig. \\ref{f1}.    Dashed red line  is for  $V_{c}= 8\\times10^{6}~cm/s$. Dotted blue line is for $V_{c}= 6\\times10^{6}~ cm/s$.} \\label{f3} \\end{figure} \\begin{figure} \\centering \\includegraphics[scale=0.55]{f3.eps} \\caption{Values  of B(G) as a function of t (years) and r(kpc) for reference values of Fig. 1.} \\label{f33} \\end{figure}"
    },
    "0910/0910.4954_arXiv.txt": {
        "abstract": "{The solar rotation profile is conical rather than cylindrical as one could expect from classical rotating fluid dynamics (e.g. Taylor-Proudman theorem). Thermal coupling to the tachocline, baroclinic effects and latitudinal transport of heat have been advocated to explain this peculiar state of rotation.} {To test the validity of thermal wind balance in the solar convection zone using helioseismic inversions for both the angular velocity and fluctuations in entropy and temperature.} {Entropy and temperature fluctuations obtained from 3-D hydrodynamical numerical simulations of the solar convection zone are compared with solar profiles obtained from helioseismic inversions.} {The temperature and entropy fluctuations in 3-D numerical simulations have smaller amplitude in the bulk of the solar convection zone than those found from seismic inversions. Seismic inversion find variations of temperature from about 1 K  at the surface up to 100 K at the base of the convection zone while in 3-D simulations they are of order 10 K throughout the convection zone up to 0.96 $R_{\\odot}$. In 3-D simulations, baroclinic effects are found to be important to tilt the isocontours of $\\Omega$ away from a cylindrical profile  in most of the convection zone helped by Reynolds and viscous stresses at some locations. By contrast the baroclinic effect inverted by helioseismology are much larger than what is required to yield the observed angular velocity profile.} {The solar convection does not appear to be in strict thermal wind balance, Reynolds stresses must play a dominant role in setting not only the equatorial acceleration but also the observed conical angular velocity profile.} ",
        "introduction": "Helioseismic data from the Global Oscillation Network Group (GONG) and the Michelson Doppler Imager (MDI) have been used to infer the rotation profile in the solar interior (e.g., Thompson et al.~1996; Schou et al.~1998). The inversion results show that isocontours of the differential rotation $\\Omega(r,\\theta)$ are conical at mid-latitude, rather than cylindrical as was expected from early numerical simulations (e.g., Glatzmaier \\& Gilman 1982; Gilman \\& Miller 1986). More recent theoretical work (Durney 1999; Kitchatinov \\& Rudiger 1995; Brun \\& Toomre 2002 (hereafter BT02); Rempel 2005; Miesch et al.~2006 (hereafter MBT06); Brun \\& Rempel 2008; Balbus et al. 2009) indicate that in order to break the Taylor-Proudman constraint of cylindrical $\\Omega$, the Sun must either have a systematic latitudinal heat transfer in its convection zone or thermal forcing from the tachocline or most likely both. This is due to the so-called thermal wind balance (Pedlosky 1987), i.e., the existence in the solar convection zone of latitudinal entropy (or temperature) variation due to baroclinic effect  can result in a rotation state that breaks the Taylor-Proudman constraint. Such latitudinal variations of the thermal properties at the solar surface have been looked for observationally by several groups since the late 60's (e.g., Dicke \\& Goldenberg 1967; Altroch \\& Canfield 1972; Koutchmy et al.~1977; Kuhn et al.~1985, 1998, Rast et al.~2008; to cite only a few). This is a difficult task since one has to correct for limb darkening effect, photospheric magnetic activity, instrument bias and many other subtle effects to extract a relatively weak signal (see Rast et al.~2008). In most cases a temperature contrast of a few degree K is found from equator to pole at the surface, the pole being warmer. In some observations a minimum at mid latitude with a warm equator and hotter polar regions is also found. The warm polar regions and cool equatorial region pattern is also found in 3-D simulation of the solar convection zone with temperature variation slightly larger (i.e., of order 10 K; BT02, MBT06). At the surface a banded structure of the temperature field (warm-cool-hot) is also found in 3-D simulation of global scale convection. While very useful and instructive, most observations are confined to the solar surface and lack the information on the deep thermal structure of the solar convection zone which is key to characterise the dynamics of the deep solar convection zone. One way to remedy that limitation is to rely on helioseismic inversions that allows us to probe deeper in the Sun and to use 3-D global simulations of the solar convection zone to guide our physical understanding.  Indeed, helioseismic inversions can give us the rotation rate, as well as the sound speed and density in the solar interior as a function of radius and latitude. Inside the convection zone the chemical composition is uniform and if we know the equation of state it is possible to determine other thermodynamic quantities like the temperature and entropy from the sound speed and density. Although, there may be some uncertainty in the equation of state, the OPAL equation of state (Rogers et al.~1996; Rogers \\& Nayafonov 2002) is quite close to the equation of state of solar material (e.g., Basu \\& Antia 1995; Basu \\& \\jcd~1997). Thus, in this work we use the OPAL equation of state to calculate the perturbations in entropy and temperature and assess how well a strict thermal wind balance is established in the solar convective envelope. To achieve this goal we make use of 2-D inversions of $\\Omega, S, T$, using the GONG and MDI data for the full solar cycle 23 and analyse our findings using 3-D simulations obtained with the ASH (anelastic spherical harmonic) code (BT02; MBT06; Miesch et al.~2008) supported by theoretical considerations on the thermal wind balance and vorticity equations.  The rest of the paper is organised as follows: in Sect.~2 we describe the data and technique used in this work while the results for the temperature and entropy inversions are described in \\S~3 along with those of 3-D simulations. In \\S~4 we discuss at length the thermal wind balance and its generalisation and interpret our seismic inversion with 3-D simulation of global scale convection. Finally, in \\S~5 we put our results in perspective and conclude.  \\begin{figure*} \\centerline{\\resizebox{0.65\\figwidth}{!}{\\includegraphics{f1a.eps} }\\hfill \\resizebox{0.65\\figwidth}{!}{\\includegraphics{f1b.eps} }\\hfill \\resizebox{0.68\\figwidth}{!}{\\includegraphics{f1c.eps} }} \\caption{The aspherical component of temperature fluctuation, $\\delta T$ obtained from the temporally averaged GONG (left panel) and MDI (middle panel) data. The right panel shows the cuts at constant latitude of $\\delta T$ obtained from MDI data along with $1\\sigma$ error estimates shown by dotted lines. All curves appear to merge at $r=R_\\odot$, because $\\delta T$ is of order of 1 K in that region.} \\label{T_obs} \\end{figure*}  \\begin{figure*} \\centerline{\\resizebox{0.65\\figwidth}{!}{\\includegraphics{f2a.eps} }\\hfill \\resizebox{0.65\\figwidth}{!}{\\includegraphics{f2b.eps} }\\hfill \\resizebox{0.65\\figwidth}{!}{\\includegraphics{f2c.eps} }} \\caption{The aspherical component of entropy fluctuation, $\\delta S$ obtained from the temporally averaged GONG (left panel) and MDI (middle panel) data. The right panel shows the cuts at constant latitude of $\\delta S$ obtained from MDI data along with $1\\sigma$ error estimates shown by dotted lines.} \\label{S_obs} \\end{figure*} ",
        "conclusions": "What can be the source of the disagreement between the inverted baroclinic contribution and the $z$ derivative of the angular velocity (i.e equations 8, or 9 and 10)? The first and easiest solution is that the inversion of the thermal quantities lack the necessary accuracy and given the increase by two orders of magnitude of the background temperature and density with depth, we end up with variations that are too large. The source of discrepancy will then be due to an overestimation of $\\delta T$ and $\\delta S$. It is not easy to decide if these inverted thermal fluctuations are too large or if the simulations (both 2-D and 3-D) underestimates the fluctuations realised in the Sun, because for instance of their limited Reynolds number. We must thus also consider the possibility that these large thermal perturbations are genuine. If this is indeed the case we need to see how we could resolve the discrepancy between the seismically inverted LHS and RHS of equation 8. As stated in section \\ref{theory}, to obtain a strict thermal wind balance as expressed in equation 8, one need to make a certain number of assumptions: adiabaticity, weak Rossby number, negligible compressibility, viscous and Reynolds stresses, stationarity. Further by considering only the hydrodynamic contributions we have omitted those associated with Maxwell stresses that are certainly present in the magnetic Sun. Let's however assume that the Maxwell stresses are not the source of the large observed discrepancy. We are confident that this is the case because we have formed temporal averages over a maximum and minimum period of activity and the differences between the two periods are about 10 times smaller that what it would required if all the sources of discrepancy were coming from the Maxwell stresses alone. We nevertheless intend to make a more systematic study of the departure of the strict thermal wind balance linked to magnetic effects (i.e. via the so called magnetic wind) by analysing the solar cycle 23 in details and by comparing with dynamo simulations of the solar convection (Brun et al.~2004). We must thus question the validity of the other hypothesis made in deriving equation 8. Clearly it is justified given the very low microscopic value of the solar kinematic viscosity to consider that the viscous terms do not contribute much. This is clearly not the case in the 3-D models where near the surface they are major contributors to the overall balance (see Figure \\ref{TW_theo}, middle panel of the bottom row), but this is due to our large effective viscosity. Assuming adiabaticity is certainly reasonable in most of the convection zone but clearly not near the surface. Since we are mostly interested in understanding the bulk dynamics of the solar convection zone, this term is indeed very small. The choice of low Rossby number that allows us to neglect $\\omega$ over $2\\Omega_0$ the planetary vorticity is certainly not justified at all scales of the turbulent velocity spectra, in particular for those scales much smaller than the Rossby radius of deformation (Pedlosky 1987). In the Sun the large range of convection scales certainly undergo different dynamics depending on how sensitive they are to the Coriolis force. The subtle angular momentum and heat redistribution realised in the  Sun is in part captured in our 3-D models. We can thus analyse if the Reynolds stresses associated with the turbulent motion indeed play a central role. As discussed in detail in Brun \\& Toomre (2002) and in \\S 4 we know that it is indeed the case in our numerical simulations (see Figure \\ref{TW_theo}, middle and right panel of the top row) even though our simulation do not possess a Reynolds number and a degree of turbulence as high as that in the Sun. We can thus expect, given the very large Reynolds number of the solar convection zone, that Reynolds stresses must play a central role in the Sun in shaping the differential rotation profile and that they somehow in part compensate the baroclinic contribution to yield the observed profile of angular velocity. This is a very interesting results, since it indicates that the differential rotation is of dynamical origin for its equatorial acceleration (as revealed by studying angular momentum transport in our simulations as in BT02 or Miesch et al.~2008) but also most certainly for its shape  with Reynolds stresses helping or opposing in some regions the baroclinic effects to break Taylor-Proudman constraint. Of course this conclusion only holds if the inverted large thermal fluctuations are real."
    },
    "0910/0910.5025_arXiv.txt": {
        "abstract": "The most X-ray luminous cluster known, RXJ1347-1145 ($z=0.45$), has been the object of extensive study across the electromagnetic spectrum.  We have imaged the Sunyaev-Zel'dovich Effect (SZE) at 90 GHz ($\\lambda = 3.3$ mm) in RXJ1347-1145 at $10''$ resolution with the 64-pixel MUSTANG bolometer array on the Green Bank Telescope (GBT), confirming a previously reported strong, localized enhancement of the SZE $20''$ to the South-East of the center of X-ray emission. This enhancement of the SZE has been interpreted as shock-heated ($> 20 \\, {\\rm keV}$) gas caused by an ongoing major (low mass-ratio) merger event. Our data support this interpretation. We also detect a pronounced asymmetry in the projected cluster pressure profile, with the pressure just east of the cluster core $\\sim 1.6 \\times$ higher than just to the west.  This is the highest resolution image of the SZE made to date. ",
        "introduction": "The rich cluster RXJ1347-1145 ($z=0.45$) is the most X-ray luminous galaxy cluster known \\citep{schindler95,schindler97,allen02} and has been the object of extensive study at radio, millimeter, submillimeter, optical and X-ray wavelengths \\citep{kitayama04,komatsu01,Gitti07,allen02,schindler97,pointecouteau99,ota08, cohen02,bradac08,miranda08}.  Discovered in the ROSAT All-Sky Survey, RXJ1347-1145 was originally thought to be a dynamically old, relaxed system \\citep{schindler95,schindler97} based on its smooth, strongly-peaked X-ray morphology--- a prototypical relaxed ``cooling-flow'' cluster. The NOBA 7 bolometer system on the 45-meter Nobeyama telescope \\citep{kitayama04,komatsu01} has made high-resolution observations ($13''$ FWHM, smoothed to $\\sim 19''$ in the presented map) of the Sunyaev-Zel'dovich effect (SZE) at 150 GHz which indicate a strong enhancement of the SZ effect $20'' \\, (170 \\, {\\rm kpc})$ to the south-east of the peak of the X-ray emission, however. Hints of this asymmetry had been seen in earlier, lower resolution measurements with the Diabolo $2.1$~mm photometer on the IRAM 30-m \\citep{pointecouteau99}.  The enhancement has been interpreted as being due to hot ($T_e > 20 \\, {\\rm keV}$) gas which is more difficult to detect using X-rays than cooler gas is, owing to the lower responsivities of imaging X-ray telescopes such as Chandra and XMM at energies above $\\sim 10 \\, {\\rm keV}$. In contrast, the SZE intensity is proportional to $T_e$ up to arbitrarily high temperatures, aside from relativistic corrections which are weak at 90 GHz, so such hot gas stands out.  The feature is consistent with the presence of a large substructure of gas in the intra-cluster medium (ICM) shock-heated by a merger, as is seen in the ``Bullet Cluster'' 1E0657-56 \\citep{markevitch02}; this interpretaion has been supported by more recent observations \\citep[e.g.][]{allen02,ota08}. {\\it Thus, rather than being an example of a hydrostatic, relaxed system, high-resolution SZE observations suggest that the observed properties of the ICM in RXJ1347-1145 are strongly affected by an ongoing merger.}  This is a striking cautionary tale for ongoing blind SZE surveys \\citep{carlstrom02}, for which useful X-ray data will be difficult or impossible to obtain for many high-$z$ systems, as well as a sign that our current understanding of nearby, well-studied X-ray clusters may be dramatically incomplete.  Reports \\citep{komatsu01,pointecouteau01,kitayama04} of a strong enhancement of the SZE away from the cluster center are based on relatively low-resolution images compared to the size of the offsets and features involved.  SZE images at lower frequencies also show show substantial offsets between the peak of X-ray and SZE emission; for instance, the 21 GHz \\citep{komatsu01} and SZE peak is $\\sim 20''$ to the SE of the X-ray peak, and the 30 GHz \\citep{reese02} SZE peak is $\\sim 13''$ to the SE of the X-ray peak.  The situation is further complicated by the presence of a radio source in the center of the cluster.  We have sought to test these claims, and to begin to untangle the astrophysics of this interesting system, with higher resolution imaging at a complementary frequency.  In this paper we present the highest angular resolution image of the SZE yet made. We observed RXJ1347-1145 with the MUSTANG 90 GHz bolometer array on the Robert C. Byrd Green Bank Telescope (GBT).  At the redshift of the cluster (and assuming $\\Omega_{\\Lambda}= 0.3,\\Omega_{tot}=1, h=0.73$) the GBT+MUSTANG $9''$ beam corresponds to a projected length of 54 kpc.  The observations are described in \\S~\\ref{sec:obs} and the data reduction in \\S~\\ref{sec:reduc}. Our interpretation and conclusions are presented in \\S~\\ref{sec:concl}. ",
        "conclusions": "\\label{sec:concl}  \\subsection{Comparison with Previous SZE Observations} Figure~\\ref{fig:Noba} presents a direct comparison of the MUSTANG and NOBA results in units of main-beam averaged Compton y parameter. For a more accurate comparison, we downgrade the resolution and pixelscale of the MUSTANG map to match that of NOBA (13'' FWHM on a 5'' pixel grid). The overall agreement between the maps is excellent, in particular as regards the amplitude and morphology of the local enhancement of the SZE south-east of the cluster core.  The largest discrepancy is south west of the cluster, where NOBA shows a $3\\sigma$ compact decrement which is absent from the MUSTANG data.  Considering the low and uniform X-ray surface brightness in the vicinity of this discrepancy (see Figure~\\ref{fig:composite}) and the higher angular resolution and lower noise of the MUSTANG data, it is likely that this feature is a spurious artifact in the NOBA map.  Both datasets also show a ridge extending north from the shock front on the eastern side of the cluster. In the 150 GHz map the feature is of marginal significance ($1-2\\sigma$); interestingly, it is clearly visible in the 350 GHz SZE increment map but K04 dismiss it due to the possibility of confusing dust emission from the nearby galaxies.  \\subsection{Empirical Model of the SZE in RXJ1347-1145} We construct a simple empirical model for the cluster SZE assuming the isothermal $\\beta$-model of \\citet{schindler97} normalized by the SZE measurement of \\citet{reese02} and \\citet{kitayama04} to describe the bulk cluster emission. We add a 5 mJy point source in the cluster core, coincident with the peak of the $\\beta$-model, and two Gaussian components in integrated pressure, one south-east and one almost directly east of the cluster center. In comparing to our 90 GHz data, we use the relatavistic correction of \\citet{sazonov98}, assuming $kT = 25 \\, {\\rm keV}$ (which reduces the amplitude of the decrement by 15\\%) for the Gaussian components and $kT = 10 \\, {\\rm keV}$ for the bulk component.  The parameters chosen (two Gaussian widths for each component, a position, a peak surface brightness, and a position angle) are shown in Table~\\ref{tbl:szmodel}.  The resulting sky image is convolved with our PSF (\\S~\\ref{sec:beam}) and transfer function (\\S~\\ref{sec:sims}).  We find that this provides a good match to the data (Figure~\\ref{fig:szmodel}). The peak comptonization at $10''$ Gaussian resolution is $3.9 \\times 10^{-4}$ on the eastern ridge and $6.0 \\times 10^{-4}$ on the region identified as a shock by Komatsu et al. When convolved to $19''$ FWHM (NOBA) resolution, we find $\\Delta y = 3.9 \\times 10^{-4}$, close to their observed value $\\Delta y = 4.1 \\times 10^{-4}$. The intent of this static, phenomenological model is simply to provide a description of the observed high angular-resolution SZE and a direct comparison of NOBA and MUSTANG results. Work is underway which will allow quantitatively determining the best fit physical model by simulataneously fitting datasets at multiple wavelengths using a Monte-Carlo Markov Chain.  This work is beyond the scope of this paper and will be presented in a follow up publication. \\begin{table*} \\begin{center} \\begin{tabular}{llll} Component & Amplitude & Offset  & Notes \\\\ &      [$y/10^{-3}$]                    &     [$''$]          &       \\\\ \\hline $\\beta$-model & $1.0$ & $0,0$ & $\\theta_c = 10''$, $\\beta=0.60$ \\\\ Shock & $1.6$ & -14, 14 & $\\sigma_1=8''$, $\\sigma_2=2''$, ${\\rm P.A.}=45^{\\circ}$ \\\\ Ridge & $1.0$ & 10, 14 & $\\sigma_1=8''$, $\\sigma_2=2''$, ${\\rm P.A.}=-15^{\\circ}$ \\\\ \\hline \\end{tabular} \\caption{Note: Offset is  (north,east) of peak X-ray position} \\label{tbl:szmodel} \\end{center} \\end{table*}   \\subsection{Multi-wavelength Phenomenology}  Our data show an SZE decrement with an overall significance of $5.4 \\sigma$.  At the center of the cluster, coincident with the peak of X-ray emission and the brightest cluster galaxy (BCG), there is an unresolved $5 \\mJy$ radio source. This flux density is consistent with the 90 GHz flux density presented in \\citet{pointecouteau01}, as well as what is expected from a power law extrapolation of $1.4$ GHz and 30 GHz measurements \\citep{NVSS,reese02}.  A strong, localized SZE decrement can be seen $20''$ to the south-east of the center of X-ray emission and clearly separated from the cluster center. Our data also indicate a high-pressure ridge immediately to the east of the cluster center.     K04 tentatively attribute the south-east enhancement to a substructure of gas $240 \\pm 183 \\, {\\rm kpc}$ in length along the line of sight, at a density (assumed uniform) of $(1.4 \\pm 0.59) \\times 10^{-2} \\, {\\rm cm^{-3}}$ and with a temperature $T_e = 28.5 \\pm 7.3 \\, {\\rm keV}$. Recent X-ray spectral measurements \\citep{ota08} with SUZAKU also indicate the presence of hot gas in the south-east region ($T_e = 25.1^{+6.1}_{-4.5} \\ ^{+6.9}_{-9.5} \\, {\\rm keV}$ with statistical and systematic errors, respectively, at 90\\% confidence level). \\citet{allen02} have reported that the slight enhancement of softer X-ray emission in this region seen by Chandra is consistent with the presence of a small substructure of hot, shocked gas. \\citet{kitayama04} attribute the hot gas to an ongoing merger in the plane of the sky. The merger hypothesis is supported by optical data, in particular, the presence of a second massive elliptical $\\sim 20''$ directly to the east of the BCG that coincides with the center of X-ray emission (and with the radio point source).  Furthermore the density and temperature of the hot substructure indicate that it is substantially overpressured compared to the surrounding ICM. Assuming a sound speed of $1600 \\, {\\rm km/sec}$ this overpressure region should relax into the surrounding ICM on a timescale $\\sim 0.1 \\, {\\rm Gyr}$, again arguing for an ongoing merger.     Our data support this merger scenario.  To put them in context, Figure~\\ref{fig:composite} shows a composite image with archival Chandra and HST data, and the weak + strong lensing mass map of \\citet{bradac08}.  We propose that the data are best explained by a merger occuring in or near the plane of the plane of the sky.  The left-hand (``B'') cluster, having fallen in from the south-{\\it west}, has just passed closest approach and is hooking around to the north-west. As the clusters merge shock forms, heating the gas in the wake of its passage. As argued by \\cite{kitayama04}, and seen in simulations \\citep{takizawa99}, the clusters must have masses within a factor of 2 or 3 of equality and a substantial ($\\sim 4000 \\, {\\rm km/sec}$) relative velocity in order to produce the high observed plasma temperatures, $T_e > 20 \\, {\\rm keV}$.  This merger geometry is consistent with the lack of structure in the line-of-sight cluster member galaxies' velocities \\citep{cohen02}.  The primary (right-hand, ``A'') cluster contains significant cold and cooling gas in its core (a ``cooling flow''). Such gas is seen to be quite robust in simulated major cluster mergers \\citep{gomez02,poole08}. Even in cases where the cooling flow is finally disrupted by the encounter, \\cite{gomez02} find a delay of $1-2$ GYr between the initial core encounter and the collapse of the cooling flow. The existence of a strong cooling flow, therefore, does not argue against a major merger in this case.  More detailed simulations could shed further light on this interesting system.  \\subsection{Broad Implications}  Since calibrating SZ observable - mass relationships is vital to understand the implications of ongoing SZE surveys, it is important to understand the mechanism by which a substantial portion of the ICM can be heated so dramatically, and how this energy is distributed through the ICM over time. Observations of cold fronts in other clusters \\citep[e.g.][]{v01} have shown that energy transport processes in the ICM are substantially inhibited, perhaps by magnetic fields. It is distinctly possible, then, that once heated by shocks, very hot phases would persist.  We have estimated the magnitude of the bias in an arcminute-resolution Compton $y$ measurement that is introduced by a hot gas phase by convolving the two gaussian components of the SZE model in Table~\\ref{tbl:szmodel} with a $1'$ FWHM Gaussian beam, typical of SZ survey telescopes such as ACT \\citep{actref,actclus} or SPT \\citep{sptref,sptcat}. Compared with the bulk emission component, also convolved with a $1'$ beam, the small-scale features are a 10\\% effect. While relatively modest this is a systematic bias in the Compton $y$ parameter which, if not properly accounted for, would result in a $20\\%$ overestimate in distances (underestimate in $H_0$) derived from a comparison of the SZE and X-ray data which did not allow for the presence of the hot gas component. To assess the impact on the {\\it scatter} in $M-y$, a larger sample of high-resolution SZE measurements is needed. A full calculation would also need to take into consideration effects such as detection apertures and the spatial filtering due to imaging algorithms, some of which would increase the importance of the effect and some of which would decrease its importance.   This is the one of a very few clusters that has been observed at sub-arcminute resolution in the SZE \\citep[see also][]{nord09}, so it is possible that many clusters exhibit similar behavior.  Such events, if their enhancement of the SZE brightness is transient, could also bias surveys towards detecting kinematically disturbed systems near the survey detection limits.  The astrophysics that has been revealed by high resolution X-ray observations, and is beginning to be revealed by high resolution SZE data, is interesting in its own right.  The SZE observations require large-aperture millimeter telescopes which have henceforth been lacking, but with both large single dishes and ALMA coming online, exciting observations will be forthcoming.  There is substantial room for improvement: since the observations we report here the GBT surface has improved from $320 \\micron$ RMS to $250 \\micron$ RMS, which will yield more than a factor of $1.5$ improvement in sensitivity. The array used in these observations, while state of the art, has not yet achieved sky photon noise limited performance; further progress is being made in this direction. Considering these facts, and that the results presented here were acquired in a short period of allocated telescope time (8h), this new high-resolution probe of the ICM has a bright future.        The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc. We thank Eiichiro Komatsu for providing the NOBA SZ map; Marusa Bradac for providing her total mass map; Masao Sako, Ming Sun, Maxim Markevitch, Tony Mroczkowski and Erik Reese for helpful discussions; and Rachel Rosen and an anonymous referee for comments on the manuscript."
    },
    "0910/0910.2433_arXiv.txt": {
        "abstract": "We have performed three searches for high-frequency signals in the solar neutrino flux measured by the Sudbury Neutrino Observatory (SNO), motivated by the possibility that solar $g$-mode oscillations could affect the production or propagation of solar $^8$B neutrinos.  The first search looked for any significant peak in the frequency range 1/day to 144/day, with a sensitivity to sinusoidal signals with amplitudes of 12\\% or greater.  The second search focused on regions in which $g$-mode signals have been claimed by experiments aboard the SoHO satellite, and was sensitive to signals with amplitudes of 10\\% or greater.  The third search looked for extra power across the entire frequency band.  No statistically significant signal was detected in any of the three searches. ",
        "introduction": "\\label{sec:intro}  Neutrinos are the only way known to directly probe the dynamics of the solar core~\\citep{Bahcall:1988}, and, through the Mikheev-Smirnov-Wolfenstein (MSW) effect~\\citep{Mikheev:1986, Wolfenstein:1977}, they can even carry information about the rest of the solar envelope.  To date, however, converting measurements of solar neutrino fluxes into constraints on solar models has proven to be difficult~\\citep{nussm}, because of the large number of co-varying parameters upon which such models are built.  A relatively simple signal that could tell us something new about the Sun would be a time-variation in the neutrino fluxes.  Over the past forty years, measurements made by solar neutrino experiments have therefore been the focus of many studies, ranging from attempted correlations with the solar sunspot cycle to open searches for signals with periods of weeks or months~\\citep{sturrock:03, sturrock:04, sturrock:05, sk:periodicity, sno:periodicity}. The shortest period examined to date is roughly one day, where the MSW effect predicts that neutrinos propagating through the Earth's core during the night will undergo flavor transformation in much the same way they do in the Sun, resulting in a net gain in the flux of electron neutrinos ($\\nu_e$s).   Although there have been occasional claims of signals on timescales similar to known variations in the solar magnetic field, in all cases there have been conflicting measurements that show the signals to be spurious or absent entirely.  We present in this article the results of a search in a new frequency regime for solar time variations. Our focus has been on signals whose periods range from 24 hours down to 10 minutes.  The motivation for such a high frequency search is in part the expectation for solar helioseismological variations on scales of order an hour or less, in particular solar `gravity modes' ($g$-modes)~\\citep{christensen}.  These $g$-modes are non-radial oscillations that are predicted to be confined to the solar core, and thus could in principle affect either neutrino production or neutrino propagation.  The neutrinos that SNO detects, those from $^8$B decay within the Sun, are particularly well-suited for our search because they are created very deep within the solar core and because their propagation is known to be sensitive to variations in the solar density profile through the MSW effect.   The effects of $g$-modes on solar neutrino fluxes have been examined by Bahcall and Kumar~\\citep{bk:gmode}, who sought to determine whether $g$-mode effects could explain the apparent solar neutrino deficit, finding that any effect was far too small to account for the roughly 60\\% discrepancy.  More recently, ~\\citet{burgess} looked at ways in which a broad spectrum of $g$-modes could alter the expectation for a solar neutrino spectral distortion caused by the MSW effect. Nevertheless, there are at this time no explicit predictions as to whether $g$-modes or any other short-timescale variations could lead to measurable solar neutrino flux variations. ",
        "conclusions": "We have performed three searches for high-frequency signals in the $^8$B solar neutrino flux, applying a Rayleigh Power technique to data from the first two phases of the Sudbury Neutrino Observatory.  Our first search looked for any significant peak in a Rayleigh Power spectrum from frequencies ranging from 1/day to 144/day.  To account for SNO's deadtime window, we calculated the expected distribution of power in each bin of the Rayleigh Power spectrum using a random walk model, thus allowing us to assign confidence levels to the observed powers.  We found no significant peaks in the data set. For this `open' peak search, we had a 90\\% probability of making a 99\\% CL detection of a signal with an amplitude of 12\\% or greater, relative to SNO's time-averaged neutrino flux.  In a second search, we narrowed our frequency band to focus on a region in which $g$-mode signals have been claimed by experiments aboard the SoHO satellite. The examined frequency range extended from 18.5/day to 19.5/day.  Again, no significant peaks in the Rayleigh Power spectrum were found, and our sensitivity for this `directed' search gave us a 90\\% probability of making a 99\\% CL detection for signals whose amplitudes were 10\\% or larger, relative to SNO's time-averaged neutrino flux.  Our third search examined the entire range of frequencies from 1/day to 144/day, looking for any evidence that additional power was present across the entire high-frequency band.   To do this, we used the distribution of frequency-specific confidence levels, determined using our random walk model.  We found that, as expected for no high-frequency variations, this distribution was flat. We showed that for a simple Gaussian white noise model, the confidence level distribution would be noticeably distorted even when the amplitudes of the contributing frequencies have an rms as small as 0.1\\%."
    },
    "0910/0910.5480_arXiv.txt": {
        "abstract": "We study the structure of neutron stars in $f(R)$ gravity theories with perturbative constraints. We derive the modified Tolman-Oppenheimer-Volkov equations and solve them for a polytropic equation of state. We investigate the resulting modifications to the masses and radii of neutron stars and show that observations of surface phenomena alone cannot break the degeneracy between altering the theory of gravity versus choosing a different equation of state of neutron-star matter. On the other hand, observations of neutron-star cooling, which depends on the density of matter at the stellar interior, can place significant constraints on the parameters of the theory. ",
        "introduction": "\\label{INTROsec} Recent interest in modified theories of gravity has been spurred by the discovery that the Universe is undergoing accelerated expansion (see, e.g.,~\\cite{P99,R98,WMAP09}). The simplest solution consistent with these observations posits a cosmological constant $\\Lambda$. The magnitude of this cosmological constant is significantly less than what was expected, and many undertakings have been made to see if there are plausible alternative explanations~\\cite{CDTT04,S08}. Outstanding questions also present themselves in the formation of singularities~\\cite{P08} and the seeming contradiction between quantum mechanics and gravity in the context of black hole thermodynamics~\\cite{W01}. All these suggest that there may yet be much to understand about the nature of gravity at extreme-curvature scales, far removed from our everyday experience.  The two most popular approaches to modifying gravity have been the introduction of an additional scalar field (e.g.~\\cite{PR03}), or the related approach of replacing the Einstein-Hilbert action with a general function of the Ricci scalar $f(R)$ (e.g.~\\cite{CDTT04}). Within either framework the additional scalar degree of freedom can be tuned to mimic the cosmological constant, or any type of cosmological evolution at cosmological scales~\\cite{W07}.  Despite the premise of such modifications, the non-linear character of gravitational theories has proven a significant obstacle to introducing new dynamical fields to drive modifications to gravity at the cosmological scale without the same fields reemerging at widely different curvature scales. One such example is the problem of ensuring that $f(R)=R\\pm\\mu^4/R$ theories pass the current Parametrized Post-Newtonian (PPN) bounds. When the new field is dynamical, the PPN parameter $\\gamma$ is forced to a value of $1/2$, which is very far from the present experimental bound~\\cite{C03}. As a result one has to choose a function $f(R)$ only from the class which can adequately suppress the new dynamical field on solar-system scales. The chameleon mechanism~\\cite{KW04,FTB07,HS08} provides such an alternative.  In addition to the PPN constraints, instabilities related to the functional form of $f(R)$ have also been studied at length. This is especially true for the Dolgov-Kawasaki instability~\\cite{DK03}, which requires that $\\partial ^2 f/ \\partial R^2 > 0$ in order that the effective mass of the equivalent scalar degree of freedom be positive. In the strong-field regime, recent results~\\cite{KM08} suggest that this very choice may well prohibit the formation of compact objects above a curvature scale readily observed.  However, the fatal curvature singularity may be avoided by the chameleon mechanism~\\cite{BL09,UH09}.  Perhaps the source of the instabilities and consistency issues many of these models encounter is the result of treating these modifications as though they are exact. The original motivation behind introducing additional functions of the curvature was to generate a new phenomenology at a specific scale. However, many of the problems encountered by $f(R)$ gravity theories originate at curvature scales far removed from the ones under consideration.  An alternative formulation for handling corrections to General Relativity is to view the new terms as only the next to leading order terms in a larger expansion. In this context there is no reason to suspect that the new phenomenology is due to new dynamical fields. The technique for handling a field expansion of this form is well developed~\\cite{EW89} and is known as perturbative constraints or order reduction~\\cite{JLM86}.  Gravity with perturbative constraints allows us to explore alternative phenomenologies of gravity while maintaining important consistency conditions including gauge invariance, the assumption that we are approximating a fundamentally second order field theory, and the conservation of stress-energy. Maintaining such constraints while enlarging the space of possible behaviors of gravitation is the goal also of the Parametrized Post-Friedman approach~\\cite{HS07,H08,FS10}.  In previous works~\\cite{DP08, CDP09}, we have analyzed the effect of treating $f(R)$ models of gravity via perturbative constraints primarily at cosmological scales. In this paper, we examine the ramifications of modifications to gravity in the context of compact objects. We show how the method of perturbative constraints allows for a consistent phenomenology for gravity on both large (Hubble-length perturbations linear in metric variables, but strongly relativistic, $L\\sim c/H_0$) and small scales (stellar scales, non-linear in metric perturbations, and strongly relativistic, $GM/rc^2\\sim 1$.)  The layout of this work is as follows. In Section~\\ref{PCsec}, we review the equations of $f(R)$ gravity treated with perturbative constraints. In Section~\\ref{SPCsec}, we derive the modified Tolman-Oppenheimer-Volkov equations and show that the exterior solution is the Schwartzchild-de Sitter metric. In Section~\\ref{MRsec}, we demonstrate that such objects are stable and we solve numerically for their mass-radius relation for a polytropic equations of state. Finally in Section~\\ref{DIS} we discuss how we can discriminate modifications to gravity from uncertainty in the neutron star equation of state. ",
        "conclusions": "\\label{DIS}  The predicted mass-radius relation for neutron stars in $f(R)$ gravity shown above differs from that computed within general relativity. However, very similar deviations in the mass-radius relation can also be obtained within general relativity by simply changing the polytropic index of the equation of state (see~\\cite{OP09} for examples). Because the equation of state of neutron-star matter is weakly constrained by current experiments, neutron-star observables that depend only on the mass and radius of the star cannot distinguish between small differences in the equation of state versus small modifications to gravity.  In~\\cite{P08-2} it was shown that observables that depend also on the effective surface gravity of neutron stars can break, in principle, this degeneracy. In particular it was shown that the Eddington luminosity $L_E^{\\infty}$ of a bursting neutron star depends directly on its effective surface gravity as \\begin{equation} L_{\\rm E}^\\infty \\equiv \\frac{4\\pi m_{\\rm p} r_{\\rm s}} {(1+X)\\sigma_{\\rm T}} \\left[\\frac{z_{\\rm s}(z_{\\rm s}+2)}{(1+z_{\\rm s})^3}\\right] \\eta\\label{eq:obs2}\\;. \\end{equation} In this equation, $m_{\\rm p}$ is the mass of the proton, $X$ is the hydrogen mass fraction in the neutron-star atmosphere, $\\sigma_{\\rm T}$ is the Thomson scattering cross section, and \\begin{equation} z_{\\rm s} = \\left(1-\\frac{2M}{R}\\right)^{-1} -1 \\end{equation} is the gravitational redshift from the neutron star surface. The parameter $\\eta$ is the ratio of the effective surface gravity of the neutron star to that calculated in GR, i.e., \\begin{equation} \\eta\\equiv\\frac{g_{\\rm eff}}{g_{\\rm GR}} \\end{equation} with \\begin{equation} g_{\\rm eff}\\equiv \\left. \\frac{1}{2\\sqrt{A}}\\frac{d \\ln B}{dr}\\right|_{r=R} \\end{equation} and \\begin{equation} g_{\\rm GR}=\\frac{1}{2R}\\left[\\frac{z_s\\left(z_s+2\\right)}{z_s+1}\\right]\\;. \\end{equation}  \\begin{figure}[t] \\includegraphics[angle=-90,scale=0.4]{alpha_rho.eps} \\caption{The central density $\\bar{\\rho}$ of neutron stars with different masses $\\bar{M}$ as a function of the parameter $\\bar{\\alpha}$. Larger positive deviations from general relativity require larger central densities for the same neutron-star mass and, therefore, lead to shorter cooling times. On the other hand, larger negative deviations require smaller central densities and lead to longer cooling time.} \\label{alpharhocgraph} \\end{figure}  We can calculate easily the value of the parameter $\\eta$ for the $f(R) = R^2$ theory considered here. From the conservation equation~(\\ref{DT}) we can write \\begin{equation} \\label{geff} g_{eff}=-\\frac{1}{\\sqrt{A}}\\frac{P'}{\\left(\\rho+P\\right)}\\;. \\end{equation} We can then evaluate the hydrostatic equilibrium equation~(\\ref{HYDRO}) to first order in $\\alpha$ by noting that \\begin{equation} \\frac{R^{(0)'}}{A^{(0)}} = -8\\pi\\left(\\frac{\\partial \\rho_0}{\\partial P_0}-3\\right)\\frac{\\left(\\rho_0 + P_0\\right)}{r^2}\\left(M^{(0)}+4\\pi P_0 r^3\\right)\\;. \\end{equation} As a result equation~(\\ref{geff}) becomes \\begin{eqnarray} g_{eff} &=& \\frac{ \\sqrt{A^{(1)}} }{r^2} \\left(M^{(1)}+4\\pi P_1 r^3\\right) -\\alpha \\left\\{8\\pi \\left(\\rho_0+P_0\\right)\\sqrt{A^{(0)}} r \\right.  \\nonumber \\\\ && \\left[ 2\\pi \\left(\\rho_0-3P_0\\right) + \\frac{2}{r^3}\\left(3-\\frac{\\partial \\rho_0}{\\partial P_0}\\right)\\left(M^{(0)}+4\\pi r^3 P_0\\right) \\right. \\nonumber \\\\ && \\left. \\left. \\left.  +\\frac{A^{(0)}}{r^4}\\left(3-\\frac{\\partial \\rho_0}{\\partial P_0}\\right) \\left(M^{(0)} +4\\pi r^3 P_0\\right)^2 \\right]\\right\\}\\right|_{r=R} \\end{eqnarray} At the surface layer of the neutron star $\\rho = P = 0$ and hence \\begin{equation} g_{eff} = \\frac{\\sqrt{A^{(1)}}}{R^2}M^{(1)}\\;. \\end{equation} Which has the same dependence on mass and radius as $g_{GR}$ does. As a result measuring $\\eta$ alone will not suffice to break the degeneracy due to the equation of state.  Nevertheless constraining observationally the cooling rates of neutron stars can offer a discriminant. A neutron star cools both through photon and neutrino emission. The photon luminosity is determined by the temperature at the photosphere, which in turn depends on the density of the photosphere. However neutrino cooling, which depends more sensitively on temperature than photon cooling does, becomes dominant for neutron stars with temperatures above $10^{10}K$, and indeed is the primary mechanism of cooling for young neutron stars (see~\\cite{PGW06} for a detailed review). The high temperature and low interaction rate make neutrino cooling particularly sensitive to the central density of the neutron star. Figure~\\ref{alpharhocgraph} shows the relation between the parameter $\\bar{\\alpha}$ and the central density of a neutron star, for three different values of the mass $\\bar{M}=0.15$, $0.125$, and $0.1$. Large positive deviations from general relativity, as measured by the parameter $\\bar{\\alpha}$ require larger central densities for neutron stars of a given mass, whereas the opposite is true for large negative deviations. As a result, because the cooling timescale scales with central density, observations of the surface temperatures of young neutron star can lead to useful constraints on the deviations from general relativity in an $f(R)$ gravity model, especially if the neutron-star masses are known.  We will study the constraints imposed on $f(R)$ gravity by current measurements of cooling rates of neutron stars in our galaxy in a forthcoming paper."
    },
    "0910/0910.5449_arXiv.txt": {
        "abstract": "The next generation of telescopes will acquire terabytes of image data on a nightly basis. Collectively, these large images will contain billions of interesting objects, which astronomers call \\textit{sources}. The astronomers' task is to construct a catalog detailing the coordinates and other properties of the sources. The source catalog is the primary data product for most telescopes and is an important input for testing new astrophysical theories, but to construct the catalog one must first detect the sources. Existing algorithms for catalog creation are effective at detecting sources, but do not have rigorous statistical error control. At the same time, there are several multiple testing procedures that provide rigorous error control, but they are not designed to detect sources that are aggregated over several pixels. In this paper, we propose a technique that does both, by providing rigorous statistical error control on the aggregate objects themselves rather than the pixels. We demonstrate the effectiveness of this approach on data from the Chandra X-ray Observatory Satellite. Our technique effectively controls the rate of false sources, yet still detects almost all of the sources detected by procedures that do not have such rigorous error control and have the advantage of additional data in the form of follow up observations, which will not be available for upcoming large telescopes. In fact, we even detect a new source that was missed by previous studies. The statistical methods developed in this paper can be extended to problems beyond Astronomy, as we will illustrate with an example from Neuroimaging.\\\\ \\textbf{Keywords:} Blind Source Detection, Multiple Testing, Astrostatistics ",
        "introduction": "The typical astronomical image records the intensity of light, over some range of frequencies, across a section of sky that contains many celestial objects of various size, shape, and luminosity. The image's pixels correspond to an array of light-sensitive detectors in the telescope, % and each pixel essentially counts how many photons have struck the corresponding detector during the exposure. But the photons recorded in the image do not come solely from the objects of interest, or \\emph{sources}; thermal noise and background emissions (collectively called \\emph{background}) corrupt the data and obscure the signature of the objects. Moreover, diffraction and atmospheric effects blur the image, reducing resolution and washing out the fainter signals. In the (astronomical) source detection problem, one is given such an image and seeks to construct a \\emph{catalog} that gives the coordinates (and often other properties) of sources in the image.  A source catalog is the basic data product of most astronomical surveys and the basic input to the scientific process. This has been true for some time. Early catalogs -- from the data of ancient astronomers Shi Shen and Hipparchus, each cataloging about 1000 stars, to the compendium of deep-sky objects produced by William and Caroline Herschel in the 1700s \\citep{catalogue-of-nebulae} % -- were based on direct visual observations. Later work, especially in the 20th century, used photographic plates, both improving resolution and allowing the detection of much fainter objects. But either way, compiling a source catalog would be a slow and painstaking affair, often requiring years to collect data on only a handful of objects. Until recently, catalogs comprising a few hundred objects were large, a few thousand were epic.  All this changed with the advent of new technologies -- digital imaging, advanced designs for telescope mirrors, and computer automation -- and with increases in available computing power and storage. With relative suddenness, astronomers found that they could observe wider, deeper, and faster than ever before. They could sweep the sky searching automatically for objects of a certain type, they could collect data on many objects in parallel, and they could observe a multitude of faint objects that would previously have gone undetected. The Sloan Digital Sky Survey \\citep{sdss} has measured hundreds of millions of objects. The upcoming Large Synoptic Survey Telescope (LSST, \\citep{lsst}) will scan the entire sky every few days, collecting several terabytes of data per night into a catalog comprising \\emph{billions} of objects. Over the past two decades, astronomy has gone from data poor to data rich.  And therein lies both opportunity and challenge. The opportunity lies in the richness and importance of the scientific questions that these massive data sets can answer. The challenge lies in the sheer scale of the data analysis. While as a general rule in science, more data is better, there is a reason that astronomers often describe the coming bounty of data in quasi-biblical terms -- a flood, a tsunami, an onslaught. The next generation of astronomical catalogs, including the LSST, will be so large that even simple operations -- such as a basic query of the entire catalog -- will be computationally prohibitive, and yes that does account for Moore's law.   This influx of data has motivated several statistical innovations in the field of Astronomy leading to the emergence of the sub-field, astrostatistics. New statistical innovations have had a significant impact on several important and cutting edge Astronomy problems, for a small sampling see \\citet{vandyk-2009-3}, \\citet{loredo}, \\citet{rice}, \\citet{npicmb}, and \\citet{richards}.   Besides the massive size of the data, another issue is controlling the rate of errors in the catalog. In the past, where every object in the catalog was observed manually, sources were often missed, located incorrectly, or created spuriously. Follow-up observations can often be made for all or most of the sources, reducing false positive and false negative identifications to a manageable level. But in the near future, the sheer number of objects in the catalog will preclude comprehensive follow-up observations by human beings. Scientific studies of the catalog will likely need to be based on samples or selections made by automatic, statistical criteria. But no matter how carefully these criteria are constructed, there will be objects that are misclassified. To use the resulting samples effectively for scientific inference, it will be necessary for the method to provide tunable control of error rates. Thus, new statistical and computational methods will be needed to construct and analyze the next generation of astronomical source catalogs.  In this paper, we develop a multiple-testing-based method for the source detection problem that has several advantages over existing techniques, especially for the analysis of large-scale surveys like the LSST. Although we discuss our method in the context of astronomical source detection, the method applies to a wide range of similar problems such as neuroimaging \\citep[e.g.][]{neuro} and remote sensing \\citep[e.g.][]{remote}, and we give such an example in a later section.  We assume that the input image is an $n\\times m$ array of pixels, with the value recorded at pixel $(i,j)$ denoted by $Y_{ij}$. The photons that contribute to $Y_{ij}$ arise from two components: \\emph{sources}, the emissions produced by the celestial objects of interest, and \\emph{background}, which includes thermal noise, the emissions of unresolved objects, interfering radiation sources, atmospheric emissions, and all other anomalies or artifacts. astronomical images essentially measure photon counts for which a Poisson model is appropriate \\citep{poissModel}. So for our base model, we assume that the $Y_{ij}$'s are independent with \\begin{equation} \\label{eq::simple-poisson-model} Y_{ij} \\;{\\rm distributed\\ as\\ } {\\rm Poisson}\\langle \\lambda_{1,ij} + \\lambda_{0,ij}\\rangle, \\end{equation} where $\\lambda_{1,ij} \\ge 0$, $\\lambda_{0,ij} \\ge 0$ denote the mean intensity of sources and background, respectively and $\\lambda_{1,ij}+\\lambda_{0,ij} > 0$.  The idea here is that the pixels are measuring the counts in disjoint cells of a Poisson random field across the sky.  This applies to good approximation for space-based observations like those reported in Section 2.  For ground-based observation, the image is in addition blurred by atmospheric turbulence, so the $Y_{ij}$'s are no longer strictly independent.  However, the Poisson model in Equation \\ref{eq::simple-poisson-model} still holds to reasonable approximation.  In data sets where the counts are high, the Poisson random field can be further approximated by a Gaussian random field.  In this paper, we utilize a technique originally developed for the Gaussian model and generalize it so that we can accommodate a wide variety of models.  The source detection problem is to identify which pixels contains sources and thus to separate the sources from the background. If $\\lambda_{1,ij} > 0$, then we take pixel $(i,j)$ to be a source pixel; otherwise, it is a background pixel. So it is natural to consider this as a multiple testing problem with the null hypothesis at each pixel being that $\\lambda_{1,ij} = 0$. At a coarse level, we want to characterize the set $\\cS = \\Set{(i,j):\\; \\lambda_{1,ij} > 0}$ of source pixels, but our more specific goal is to identify and locate the underlying sources, so that an accurate catalog can be constructed. This requires a more stringent criterion for success because the objects are coherent, localized aggregates. As Figure \\ref{fig:type1and2} shows, with the same number of pixel-wise type I and type II errors, it is possible to get widely varying accuracy in the resulting catalog. Put another way, our loss function operates on the catalog, not the pixels themselves.  \\begin{figure}[h!] \\centerline{\\includegraphics[scale=.35]{typeIandII.ps}} \\caption{The 45 pixels that have any overlap with the red circle are considered sources and black pixels indicate sources detected via some detection algorithm. Each of the three images have 6 Type I error pixels and 24 Type II error pixels. The detection in the left image captures the center of the source but clearly misses the shape. The detections in the center image show several different sources that have some overlap with the true source. The detection in the right image captures the center of the true source and has some spurious noise detections. While all these pictures have the same number of pixel-wise Type I and Type II errors they lead to very different conclusions about the number, shape and location of the sources in the image, thus a testing criteria based on the aggregate sources, instead of pixels, is necessary.   } \\label{fig:type1and2} \\end{figure}  In this paper, we extend the False Cluster Proportion (FCP) controlling procedures introduced by Perone Pacifico et al. (2004) to make it effective for controlling the rate of false sources detected in astronomical images. As we will show below, the original FCP procedure does not perform well with the Poisson statistics common in the source detection problem and even where the Gaussian assumption holds, does not yield sufficient power to be viable for the astronomical source detection problem. We generalize the technique so that it applies to a wider range of noise models. We also introduce a new transform, which we call the \\emph{Multi-scale Derivative}, that enhances sources and significantly improves power. Taken together, these extensions lead to a new procedure that we call the Multi-scale False Cluster Proportion (MSFCP) procedure. This gives a powerful source detection technique that provides rigorous control over the rate of false \\emph{sources}, where techniques in current use provide control over the rate of false \\emph{pixels}, if they provide any control at all. We demonstrate both the excellent power and error control on a very deep and high resolution telescope image from the Chandra X-Ray Observatory. MSFCP has detection power competitive with the existing procedures, even detecting a source that was overlooked in previous studies while at the same time maintaining rigorous control over the rate of false sources. Although the source detection problem is ubiquitous in Astronomy it also occurs in other settings and we present an example of how FCP concepts can be extended to detecting bands of neural activity in the brain.   Due to its prevalence and difficulty, there have been a variety of approaches to the source detection problem, both in the statistical and astronomical literature. Given the recent explosion in research on multiple testing, there is a plethora of available techniques that can be applied directly to the pixel-wise hypothesis tests to reconstruct $\\cS$, including \\citet{bh}[BH], \\citet{stepup}, \\citet{storey}, \\citet{suncai}, and \\citet{rice} . However, because these methods give error-rates in terms of the individual pixel-wise tests, it is not obvious how to translate these error-rates to make inferences about the underlying sources. Multiple testing approaches that are less pixel-centered have been developed in the related problem of analyzing functional magnetic resonance imaging (fMRI) data. In this problem, the sources are regions of neural activity that reveal themselves through a measurable change in blood flow response. \\cite{wor96} and \\citet{wor02} use level sets of a random field formed from test statistics to identify regions containing sources. \\citet{cba} construct a test based on clusters instead of pixels in the fMRI setting, but this technique takes advantage of a temporal dimension that is not available in many general (e.g., astronomical) problems.  Source detection understandably garners much attention in the astronomical literature. Astronomers address source detection as one step in a data-processing pipeline -- the series of operations performed on the data from collection until catalog. These include, but are not limited to, corrections for atmospheric effects, image registration, filtering out unwanted signals, as well as source detection. These pipelines are typically planned and developed well before the instrument is operational. During the planning stage, simulations are run to test the pipeline including the source detection algorithm \\citep[e.g.,][]{sehgal}. Astronomers typically do not require their algorithms to satisfy any formal performance criteria; rather, they apply an empirical criterion, using data simulated to look like a real telescope image to calibrate the error rate for detected sources that will be expected in practice. Of course, this depends on the simulated and real data being both quantitatively and qualitatively similar. While great effort and ingenuity are applied in constructing realistic simulations -- sometimes years of computing time for a single run -- the simulations still rely on untested and unstated assumptions that may fail when the instrument comes on line.  The source detection methods used by astronomers fall into three general classes: simple thresholding, peak-finding algorithms, and Bayesian algorithms. Simple thresholding consists of choosing an intensity cutoff for pixel-wise statistics and classifying any pixel above threshold as a source. It is popular because it simple and fast, easily computed by the popular SExtractor software \\citep{sext}. Before thresholding, filters are often applied to the raw data to suppress confounding background signals. In \\citet{vik} and \\citet{melin}, Matched Filters are used to isolate the signal, and then thresholds are determined using simulated data to create catalogs in X-ray and Radio telescope images respectively. \\citet{2002AJ....123.1086H} also use simple thresholding, choosing the threshold to control the False Discovery Rate via the method of \\citet{bh}, but they use simulated telescope images to calibrate between rate of false pixels and the rate of false sources.  Peak-finding algorithms search for local maxima in the denoised image and catalog them as sources. \\citet{vale} and \\citet{sehgal} use peak-finding algorithms to look for large galaxy clusters in simulated radio telescope images. The wavelet-based technique of \\citet{peter} is commonly used to detect X-ray sources as in \\citet{valt} and \\citet{1msec}. \\citet{Damiani} and \\citet{gonzaleznuevo-2006-369} provide a good overview of the popular Mexican Hat wavelet as a tool for source detection. Several implementations exist in software and most are specifically tailored for a certain instrument or type of problem. Pixel-wise error rates for wavelet source detection can sometimes be determined analytically but are more often approximated from simulations. As with simple thresholding, simulations are often used to indirectly estimate the rate of false sources, the error rate of interest, from the rate of false pixels.  Bayesian techniques have become popular with astronomers and several have applied Bayesian methods to source detection as in \\citet{2007ApJ...661.1339S}, \\citet{strong}, and \\citet{2003MNRAS.338..765H}. Bayesian detection algorithms typically define models for sources, often two-dimensional Gaussians, and attempt to distinguish them from backgrounds via inference on a posterior distribution. \\citet{gug} use Bayesian mixture models to separate sources from background, while \\citet{fastBayes} propose ways to speed up Bayesian source detection, which can be computationally slow for large problems. Bayesian algorithms typically require more assumptions to be made a priori and can be more computationally expensive than thresholding or peak-finding algorithms.   Our goal in source detection is to detect the relevant objects, not pixels, while controlling the error rates. Statisticians have developed numerous methods for dealing with error rates but have not been focused on detecting aggregate objects, while astronomers have been thinking about detecting objects, but without a rigorous approach to controlling error rates. We propose a technique that does both: control the error rates and make our inference about the sources themselves.  We demonstrate our techniques on an important data set from the Chandra X-ray Observatory, one of the most powerful telescopes in existence. For the Chandra data, our goal is to detect the X-ray sources with a bound on the error rate for sources. We describe an approach due to \\citet{Pacifico04} that gives us a probabilistic bound on the error rate for sources, and apply it to the Chandra data. We find that while this approach gives us the error control we want, it does not have good power when compared to other techniques. We introduce a generalization of the technique that allows for us to keep a probabilistic bound on the error rate for sources under more general conditions. We then introduce a new Multi-scale technique that is designed to enhance sources and thus increase power. We then integrate it with our generalized procedure to get the MSFCP procedure which increases power while maintaining control over the error rate. This improvement is evident when we revisit the Chandra data -- we show that our power using MSFCP is competitive to two algorithms commonly used by astronomers, but with superior error control. Furthermore, using our procedure we detect a X-ray source which had gone undetected in the original analysis of the data by astronomers. We then provide a brief description of how these techniques can be used outside the realm of Astronomy with an application to high-resolution neuroimaging data. We conclude with an overview of our results and directions for further study. ",
        "conclusions": "We have extended the False Cluster Proportion Multiple Testing Procedure so it can now be applied to a wide variety of problem in Astronomy as well as other fields. By moving away from theoretical approximations, like Piterbarg's approximation, we can now calculate confidence supersets for a large variety of noise conditions. This allows us to create source catalogs with a probabilistic guarantee that they are not overly polluted with false detections without having to make any difficult to check assumptions about the data or the science behind the data. The Multi-scale Derivative procedure enhances the types of sources we typically see in a telescope image and we have shown that they can also enhance the power of False Cluster Proportion source detection algorithms. This is evidenced by our analysis of the CDFS data. The catalog published in GI uses two different algorithms and then has to follow up each detection, whereas we run MSFCP and can make the statement that with probability .95 less than 10\\% of our detections are false, and still get virtually the same catalog. When we allow 20\\% of our detections to be false we make two new detections, one of which we have verified is a real source that was missed in the original analysis. Our procedure provides rigorous error control, which is lacking in the current techniques used by astronomers. At the same time we have demonstrated that the detection power is comparable and in fact we have detected sources that they have missed. Controlling false detections without the need for expensive follow up observations will be critical as the next generation of telescopes will provide a deluge of data that will be impossible to process manually.  We have also shown that these ideas can be generalized to other types of problems where we are dealing with points spread along a plane instead of a regular grid of pixels in an image. We can further modify our definition of cluster to differentiate between different types of objects. We can then control the proportion of false clusters, allowing for different classes of clusters. In the future, we believe we can further extend our framework so that we can apply False Cluster Proportion techniques on a wide variety of statistical clustering problems, in which we are trying to detect clusters of data without making false detections.   {\\scriptsize"
    },
    "0910/0910.3954_arXiv.txt": {
        "abstract": "We study the dynamics of stellar-mass black holes (BH) in star clusters with particular attention to the formation of BH-BH binaries, which are interesting as sources of gravitational waves (GW). In the present study, we examine the properties of these BH-BH binaries through direct N-body simulations of star clusters using the NBODY6 code on Graphical Processing Unit (GPU) platforms. We perform simulations for star clusters with $\\leq 10^5$ low-mass stars starting from Plummer models with an initial population of BHs, varying the cluster-mass and BH-retention fraction. Additionally, we do several calculations of star clusters confined within a reflective boundary mimicking only the core of a massive star cluster which can be performed much faster than the corresponding full cluster integration. We find that stellar-mass BHs with masses $\\sim 10\\Ms$ segregate rapidly ($\\sim 100$ Myr timescale) into the cluster core and form a dense sub-cluster of BHs within typically $0.2 - 0.5$ pc radius. In such a sub-cluster, BH-BH binaries can be formed through 3-body encounters, the rate of which can become substantial in dense enough BH-cores. While most BH binaries are finally ejected from the cluster by recoils received during super-elastic encounters with the single BHs, few of them harden sufficiently so that they can merge via GW emission within the cluster. We find that for clusters with $N \\ga 5\\times 10^4$, typically 1 - 2 BH-BH mergers occur per cluster within the first $\\sim 4$ Gyr of cluster evolution. Also for each of these clusters, there are a few escaping BH binaries that can merge within a Hubble time, most of the merger times being within a few Gyr. These results indicate that intermediate-age massive clusters constitute the most important class of candidates for producing dynamical BH-BH mergers. Old globular clusters cannot contribute significantly to the present-day BH-BH merger rate since most of the mergers from them would have occurred much earlier. On the other hand, young massive clusters with ages less that 50 Myr are too young to produce significant number of BH-BH mergers. We finally discuss the detection rate of BH-BH inspirals by the ``LIGO'' and ``Advanced LIGO'' GW detectors. Our results indicate that dynamical BH-BH binaries constitute the dominant channel for BH-BH merger detection. ",
        "introduction": "Star clusters, \\eg, globular clusters (henceforth GC), young and intermediate-age massive clusters and open clusters harbor a large overdensity of compact stellar remnants compared to that in the field by virtue of their high density and the mass segregation of the compact remnants towards the cluster core \\citep{hv83}. These compact stars, which are neutron stars (henceforth NS) and black holes (henceforth BH), are produced by the stellar evolution of the most massive stars within the first $\\sim$ 50 Myr after cluster formation. These compact stars, being generally heavier than the remaining low-mass stars of the cluster, segregate quickly (within one or a few half-mass relaxation times) to the cluster core, forming a dense sub-cluster of compact stars. Such a compact-star sub-cluster is of broad interest as it efficiently produces compact-star binaries through dynamical encounters \\citep{hv83,h.et.al92}, \\eg, X-ray binaries and NS-NS and BH-BH binaries, which are of interest for a wide range of physical phenomena. For example, X-ray binaries are the primary sources of GC X-ray flux, while mergers of tight NS-NS binaries (through emission of gravitational waves) is the most likely scenario for the production of short duration GRBs and both NS-NS and BH-BH mergers are very important sources of gravitational waves (henceforth GW) detectable by future gravitational wave observatories like ``Advanced LIGO'' (AdLIGO) and ``LISA'' \\citep{as2007}. Double-NS systems are also interesting because they could be observable as double-pulsar systems \\citep{rsm2008}, allowing important tests of General Relativity \\citep{krm2008}. In the present work, we investigate the dynamics of stellar-mass BHs in star clusters, with particular emphasis on dynamically formed BH-BH binaries. Such binaries are strong sources of GWs as they spiral-in through GW radiation, a process detectable out to several thousand Mpc distances.  Tight BH binaries that can merge within a Hubble time can also be produced in the galactic field from tight stellar binaries, which result from stellar evolution of the components (\\eg, \\citealt{bul2003,osh2007,osh2008,bky2007}). However, \\citet{bky2007} have shown with revised binary-evolution models that the majority of potential BH-BH binary progenitors actually merge because of common envelope (CE) evolution which occurs when any of the binary members crosses the Hertzsprung gap. They found that this reduces the merger rate by a factor as large as $\\sim 500$ and the resulting AdLIGO detection rate of BH-BH binary mergers in the Universe from primordial binaries is only $\\sim 2$ yr$^{-1}$, subject to the uncertainties of the CE evolution model that has been incorporated. Note that the corresponding NS-NS merger rate, as the above authors predict, is considerably higher, $\\sim 20$ yr$^{-1}$. In view of the above result, the majority of merging BH binaries are possibly those that are formed dynamically in star clusters.  As studied by several authors earlier \\citep{mt2006,mak2007}, black holes, formed through stellar evolution, segregate into the cluster core within $\\sim 0.3$ pc and form a sub-cluster of BHs, where the density of BHs is large enough that BH-BH binary formation through 3-body encounters becomes important \\citep{hh2003}. These dynamically formed BH binaries then ``harden'' through repeated super-elastic encounters with the surrounding BHs \\citep{h75,bg2006}. The binding energy of the BH binaries released is carried away by the single and binary BHs involved in the encounters. This causes the BHs and the BH binaries to get ejected from the BH-core to larger radii of the cluster and they heat the cluster while sinking back to the core through dynamical friction \\citep{mak2007}. Most of the energy of the sinking BHs is deposited in the cluster core as the stellar density is much higher there. As the BH binaries harden, the encounter-driven recoil becomes stronger and finally the recoil is large enough that the encountering single BH and/or the BH binary escape from the cluster (see Sec.~\\ref{res}, also \\citealt{bcq2002}). Because of the associated mass-loss from the cluster core, this also results in cluster heating. These heating mechanisms result in an expansion of the cluster, as studied in detail by several authors, \\eg, \\citet{mt2006,mak2007}. \\citet{mak2007} found notable agreement between the core expansion as obtained from their N-body simulations with the observed age-core radius correlation for star clusters in the Magellanic Clouds.  The dynamics of stellar mass BHs and formation of BH binaries in star clusters and the resulting rate of GW emission from close enough BH binaries has been studied by several authors using N-body integrations \\citep{pzm2000}, numerical 3-body scattering experiments \\citep{gul2004} or Monte-Carlo methods \\citep{olr2006}. \\citet{pzm2000} considered the rate of mergers of escaping BH binaries from various stellar systems, \\eg, massive GCs, young populous clusters and galactic nuclei. They estimated the merger rates within 15 Gyr (their adopted age of the Universe) by assuming the binding-energy ($E_b$) distribution of the escaping BH binaries to be uniformly distributed in $\\log{E_b}$, as inferred from simulations of $N \\approx 2000$ or 4000 star clusters (see \\citealt{pzm2000} for details). Considering the space-densities of the different kinds of star clusters, they found that while for LIGO the corresponding total (\\ie, contribution from all types of star clusters) detection rate is negligible, it can be as high as $\\sim 1$ day$^{-1}$ for AdLIGO. \\citet{gul2004} performed sequential numerical integrations of BH binary-single BH close encounters in a uniform stellar background, where, in between successive encounters, the BH-binaries were evolved due to GW emission. From such simulations, these authors studied the growth of BHs through successive BH-binary mergers for the first time. In a more self-consistent study of the growth of BHs and the possibility of formation of intermediate mass black holes through successive BH-binary mergers, \\citet{olr2006} used the Monte-Carlo approach and considered ``pure'' BH clusters that are dynamically detached from their parent clusters which can be expected to form due to mass stratification instability \\citep{spz}. These authors considered dynamically formed BH binaries from 3-body encounters in such BH clusters utilizing theoretical cross-sections of 3-body binary formation (see \\citealt{olr2006} and references therein). Furthermore, they included both binary-single and binary-binary encounters. Considering the sub-set of BH binaries formed in the BH clusters that merge within a Hubble time, they determined typically a few AdLIGO detections per year for old GCs.  While such results are remarkable and promising, a more detailed study of the dynamics of stellar-mass BHs in star clusters with a realistic number of stars is essential. As the cross-sections of different processes governing the dynamics of BH binaries, \\viz, 3-body binary formation, binary-single star encounters and ionization have different dependencies on the number of stars $N$ of the cluster, an extrapolation to much different $N$ can be problematic. Hence it is important to consider clusters with values of $N$ appropriate for massive clusters or GCs. Also, exact treatments of the various dynamical processes are crucial for realistic predictions of BH-BH merger rates. Finally, a more careful study of the different dynamical processes leading to the formation and evolution of BH binaries in star clusters is needed to understand better under which conditions tight inspiralling BH binaries can be formed dynamically from a star cluster.  In the present work, we make a detailed study of the dynamics of BH-BH binaries formed in a BH sub-cluster, as introduced above. In particular, we investigate whether hard enough BH binaries that can merge via gravitational radiation in a Hubble time within the cluster or after getting ejected from the cluster, can be formed in such a sub-cluster. To that end, we perform direct N-body integrations of concentrated star clusters (half-mass radius $r_h \\leq 1$ pc) consisting of $N \\leq 10^5$ low-mass stars in which a certain number of stellar-mass BHs is added, representing a star cluster with an evolved stellar population.  The present paper is organized as follows: In Sec.~\\ref{sim} we describe our simulations in detail. We discuss the various elements and assumptions of the simulations and summarize all the runs that we perform (Sec.~\\ref{runs}). We also discuss the use of a reflective boundary to simulate only the core of a star cluster (Sec.~\\ref{reflct}). In Sec.~\\ref{res} we discuss our results with particular emphasis on the dynamical BH binaries formed during the simulations. We discuss the BH-BH mergers that occur within the clusters and also the merger timescales of the BH binaries that escape from the clusters (Sec.~\\ref{mrgesc}) and obtain the distributions of merger-times for both cases (Sec.~\\ref{tdist}). Finally, in Sec.~\\ref{discuss} we interpret our results in the context of different types of star clusters that are observed and provide estimates of BH-BH merger detection rates. ",
        "conclusions": "\\label{discuss}  Our present study indicates that centrally concentrated star clusters, with $N \\ga 4.5\\times 10^4$ are capable of dynamically producing BH binaries that can merge within a few Gyr, provided a significant number of BHs are retained in the clusters after their birth. The results of our simulations (see Table~\\ref{tab1}) imply that most of the BH-BH mergers occur within the first few Gyr of cluster evolution for both mergers within the cluster and mergers of escaped BH binaries.  \\begin{figure} \\includegraphics[width=8.5cm,angle=0]{f5.eps} \\includegraphics[width=8.5cm,angle=0]{f6.eps} \\caption{{\\bf Top:} Distribution of the merger times $t_{mrg}$ for BH binary mergers within the cluster for the models of Table~\\ref{tab1}. {\\bf Bottom:} Distribution of the merger times $t_{mrg}$ for escaped BH binaries for the models of Table~\\ref{tab1} (see text).} \\label{fig:mrgdist} \\end{figure}  The above results imply that an important class of candidates for dynamically forming BH binaries that merge at the present epoch are star clusters with initial mass $M_{cl} \\ga 3\\times 10^4 \\Ms$\\footnote{To correlate with the observed clusters, we consider the mass $M_{cl}$ of the parent cluster, \\ie, the cluster mass before the BHs are formed through stellar evolution, which are heavier than the clusters of low mass stars that we model, for the same value of $N$. For a given simulated cluster, we estimate the corresponding parent mass $M_{cl}$ by weighting the mean stellar mass of $\\langle m \\rangle_{cl}\\approx 0.6 \\Ms$ of a Kroupa IMF with star from $0.1\\Ms$ upto $100\\Ms$ with the total number of stars $N$ for that cluster.}, which are less than few Gyr old. Such clusters represent intermediate-age massive clusters (hereafter IMC) with initial masses close to the upper-limit of the initial cluster mass function (ICMF) in spiral \\citep{wkl2004,lrs2009a} and starburst \\citep{gls2006} galaxies. While it is not impossible to obtain BH-BH mergers within a Hubble time from lower-mass clusters, the overall BH-BH merger and escape rates strongly decrease with cluster-mass as Table~\\ref{tab1} indicates. For $M_{cl} \\la 1.5 \\times 10^4 \\Ms$, mergers already become much rarer (see Table~\\ref{tab1}). Due to the statistical nature of merger or ejection events it is ambiguous to set any well-defined limit on the cluster-mass beyond which these events become appreciable (such an estimate would also require a much larger number of N-body integrations). In view of our results, $M_{cl} \\approx 3\\times 10^4 \\Ms$ is a representative lower limit beyond which an appreciable number of mergers and escapers merging within a Hubble time can be obtained.  Old globular clusters, which can be about 10 times or more massive, are expected to produce mergers or escapers more efficiently. As the timescale of depletion of BHs from the BH cluster is nearly independent of the parent cluster mass (see Sec.~\\ref{tdist}), GCs can also be expected to produce BH-BH mergers over similar time-span as the IMCs, \\ie, within the first few Gyr of evolution. Since GCs are typically much older ($\\sim 10 {\\rm ~Gyr}$), they do not contribute significantly to the present-day merger rate, since most of the mergers from them would have occurred earlier. Considering the light-travel time of $\\approx 4.5 {\\rm ~Gyr}$ from the maximum distance $D\\approx 1500 {\\rm ~Mpc}$ form which these BH-BH binaries can be detected by ``AdLIGO''(see below), only GCs close to the above distance could contribute detectable events, mostly from escaped BH-BH binaries.  On the other hand, young massive clusters with ages less than 50 Myr, representing star clusters near the high-mass end of the ICMF \\citep{lrs2009b}, are generally too young to produce BH-BH mergers. All models in Table~\\ref{tab1} produce mergers significantly later than this age (except one escaped BH binary in each of the models C65K110 and C100K200). Hence, IMCs seem to be most likely candidates for producing observable BH-BH mergers dynamically.  \\subsection{Detection rate}\\label{rate}  We now make an estimate of the BH-BH merger detection rate from IMCs by ground-based GW observatories like LIGO and AdLIGO. In estimating the overall BH-BH merger rate using the results of our model clusters, one needs to consider the distribution of the cluster parameters that are varied over the models, \\viz, cluster mass, half-mass radius, and BH retention fraction. Such distributions are far from being well determined, except for the mass distribution for young clusters in spiral and starburst galaxies \\citep{bik2003,blm2003,gls2004}. Therefore, determination of an overall merger rate considering the distribution of our computed clusters can be ambiguous. Hence, as a useful alternative, we determine the BH binary merger detection rates for each of the cluster models in Table~\\ref{tab1} that gives an appreciable number of mergers, for a representative density of IMCs. Such an approach has been considered by earlier authors, \\eg, \\citet{olr2006} and can provide a reasonable idea of the rate of detection of BH-BH mergers from IMCs.  As an estimate of the space density of IMCs, we adopt that for young populous clusters in \\citet{pzm2000}, which has a similar mass-range as the IMCs: \\begin{equation} \\rho_{cl} = 3.5 {\\rm ~} h^3 {\\rm ~Mpc}^{-3}, \\label{eq:rypc} \\end{equation} where $h$ is the Hubble parameter, defined as $H_0/100{\\rm ~km~s}^{-1}$, $H_0$ being the Hubble constant \\citep{pbl93}. The above space density has been derived from the space densities of spiral, blue elliptical and starburst galaxies \\citep{hey97} assuming that young populous clusters have the same specific frequencies ($S_{N}$) as old GCs \\citep{vdb95,mcl99}, but in absence of any firm determination of the $S_{N}$s of the former. We compute the detection rate for each model cluster assuming that it has a space density of the above value.  The LIGO/AdLIGO detection rate of BH-BH mergers from a particular model cluster can be estimated from (\\citealt{bky2007} and references therein) \\begin{equation} {\\mathcal R}_{LIGO} = \\frac{4}{3}\\pi D^3\\rho_{cl}{\\mathcal R}_{mrg}, \\label{eq:ligorate} \\end{equation} where ${\\mathcal R}_{mrg}$ is the compact binary merger rate from a cluster and $D$ is the maximum distance from which the emitted GW from a compact-binary inspiral can be detected. $D$ is given by \\begin{equation} D = D_0 \\left(\\frac{M_{ch}}{M_{ch,nsns}}\\right)^{5/6}, \\label{eq:range} \\end{equation} where $D_0=18.4$ and 300 Mpc for LIGO and AdLIGO respectively. The quantity $M_{ch}$ is the ``chirp mass'' of the compact binary with component masses $m_1$ and $m_2$, which is given by \\begin{equation} M_{ch} = \\frac{(m_1 m_2)^{3/5}}{(m_1 + m_2)^{1/5}}, \\label{eq:chirp} \\end{equation} and $M_{ch,nsns}=1.2\\Ms$ is that for a binary with two $1.4\\Ms$ neutron stars.  For a BH binary with $m_1=m_2=10\\Ms$, $M_{ch}=8.71\\Ms$ which gives $D \\approx 1500{\\rm ~Mpc}$ for AdLIGO. The AdLIGO detection rates ${\\mathcal R}_{\\rm AdLIGO}$ (mean over 3 Gyr taking into account the time of escape $t_{esc}$ of the escaped binaries) for the model clusters are shown in Table~\\ref{tab1}, where the currently accepted value of the Hubble parameter $h=0.73$ is assumed. The error in each detection rate is simply obtained from the Poisson dispersion of the total number of mergers for the corresponding cluster. Note that these detection rates are for clusters with solar-like metallicity which is implied by our assumption of $10\\Ms$ BHs for all the clusters.  To obtain a basic estimate of the overall detection rate of BH-BH mergers from IMCs, we consider the subset of our computed models that are isolated clusters with full BH retention (see Table~\\ref{tab1}). We take the mass function of the IMCs to be a power law with index $\\alpha = -2$ which is the approximate index of the ICMF in spiral and starburst galaxies (see \\eg, \\citealt{gls2004,lrs2009a}). Then the weighted average of the corresponding AdLIGO detection rates is $\\overline{\\mathcal R}_{\\rm AdLIGO} \\approx 31(\\pm 7)$ yr$^{-1}$, which estimates the total present-day detection rate of BH-BH mergers from IMCs expected for AdLIGO. The corresponding LIGO detection rate is negligible, $\\overline{\\mathcal R}_{\\rm LIGO} \\approx 7.4\\times 10^{-3}{\\rm ~yr}^{-1}$. Note that these BH-BH detection rates are only lower limits. First, the observed population of star clusters can be an underestimation by a factor of 2 owing to their dissolution in the tidal field of their host galaxies (see \\citealt{pzm2000} and references therein). Second, the above detection rates are only from IMCs and there can be additional contributions from GCs (see above).  In comparing the AdLIGO detection rates from our computations with those from earlier works, we note that our rates are typically an order of magnitude smaller than those of \\citet{pzm2000}, but about one order of magnitude larger than those obtained by \\citet{olr2006}. The principal origin of the former difference is due to the fact that while we considered only IMCs distinguishing them as the most appropriate candidates for producing present-day BH-BH mergers (see above), \\citet{pzm2000} also included GCs, which have considerably larger spatial density and also larger fraction of BH-BH binaries merging within the Hubble time, as obtained from their analytic extrapolations. Note that the number of escapers as obtained by them for young populous clusters and the fraction of them merging within a Hubble time (see Table~1 of \\citealt{pzm2000}) is similar to that obtained from the present computations, implying qualitative agreement. The above authors apparently did not consider the time scales of formation and depletion of the BH sub-systems in their preliminary study. On the other hand, although \\citet{olr2006} considered clusters significantly more massive than our's in their Monte-Carlo approach, so that larger merger detection rates can be expected, their clusters were much older (8 Gyr and 13 Gyr) than IMCs which, in accordance with our results, accounts for their much lower detection rate.  It is interesting to note that the dynamical BH-BH merger detection rates obtained by us are typically an order of magnitude higher than that from primordial stellar binaries as predicted by \\citet{bky2007} based on their revised binary evolution model, and is similar to that for the isolated NS-NS binaries derived by them. Hence our results imply that dynamical BH-BH binaries constitute the dominant contribution to the BH-BH merger detection. Thus, the dynamical BH-BH inspirals from star clusters seem to be a promising channel for GW detection by the future AdLIGO, although their estimated detection rate with the present LIGO detector is negligible, in agreement with the hitherto non-detection of GW.  \\subsection{Limitations and outlook}\\label{lim}  The work presented here is a first step towards a detailed study of the dynamics of stellar-mass BHs in star clusters and the consequences for GW-driven BH mergers, and improvements in several directions are possible. First, we do not consider the initial phase of the cluster here when BHs form through stellar evolution and a more consistent approach would be to begin with a star cluster with a full stellar spectrum and produce the BHs from stellar evolution. Note however that in the present study we have inserted numbers of BHs in our clusters of low-mass stars similar to what would have been formed from the stellar evolution of a cluster following a Kroupa IMF (see Sec.~\\ref{runs}). Stellar evolution also produces NSs (about twice in number as the BHs) which also segregate to the central region of the cluster. It is interesting to study the dynamics of the NS cluster and how it is affected by the (more concentrated) BH cluster --- in particular the formation of tight inspiralling NS-NS and NS-BH binaries, which are important for both GW detection and GRBs.  Another aspect that we do not consider in our present study is the effect of primordial stellar binaries. Tight stellar binaries aid the formation of compact binaries through double-exchanges \\citep{gr2006}, in addition to the 3-body mechanism (see Sec.~\\ref{intro}) which can increase the number BH-BH (also BH-NS and NS-NS) binaries formed and hence the merger rate. Thus the study of BH-BH binaries in star clusters with primordial binaries is an important next step. Such studies are in progress and will be presented in future papers.  Finally, the number of N-body computations in this initial study is not enough to obtain the BH-BH merger rate as a function of cluster mass and BH retention fraction with a reasonable accuracy. There are typically one integration per cluster model. To obtain merger rate dependencies with cluster parameters, \\eg, mass, half mass radius, binary fraction and BH retention, many computations are needed within smaller intervals, which involves much larger time and computing capacity than that utilized for the present project. Such results, combined with improved knowledge of cluster parameter distributions and BH formation through supernova explosions of massive stars, would provide a robust estimate of the BH-BH merger detection rate. Conversely, with such a merger rate function, the detection of BH-BH mergers by GW detectors like AdLIGO in the near future will shed light on the above mentioned long-standing questions."
    },
    "0910/0910.3486_arXiv.txt": {
        "abstract": "{Earth-sized planets around nearby stars are being detected for the first time by ground-based radial velocity and space-based transit surveys. This milestone is opening the path toward the definition of missions able to directly detect the light from these planets, with the identification of bio-signatures as one of the main objectives. In that respect, both the European Space Agency (ESA) and the National Aeronautics and Space Administration (NASA) have identified nulling interferometry as one of the most promising techniques. The ability to study distant planets will however depend on the amount of exozodiacal dust in the habitable zone of the target stars.} {We assess the impact of exozodiacal clouds on the performance of an infrared nulling interferometer in the Emma X-array configuration. The first part of the study is dedicated to the effect of the disc brightness on the number of targets that can be surveyed and studied by spectroscopy during the mission lifetime. In the second part, we address the impact of asymmetric structures in the discs such as clumps and offset which can potentially mimic the planetary signal.} {We use the \\darwinsim{} software which was designed and validated to study the performance of space-based nulling interferometers. The software has been adapted to handle images of exozodiacal discs and to compute the corresponding demodulated signal.} {For the nominal mission architecture with 2-m aperture telescopes, centrally symmetric exozodiacal dust discs about 100 times denser than the solar zodiacal cloud can be tolerated in order to survey at least 150 targets during the mission lifetime. Considering modeled resonant structures created by an Earth-like planet orbiting at 1\\,AU around a Sun-like star, we show that this tolerable dust density goes down to about 15 times the solar zodiacal density for face-on systems and decreases with the disc inclination.} {Whereas the disc brightness only affects the integration time, the presence of clumps or offset are more problematic and can seriously hamper the planet detection. The upper limits on the tolerable exozodiacal dust density derived in this paper must be considered as rather pessimistic, but still give a realistic estimation of the typical sensitivity that we will need to reach on exozodiacal discs in order to prepare the scientific programme of future Earth-like planet characterisation missions.} ",
        "introduction": "\\label{sec:intro}  The possibility of identifying habitable worlds and even biosignatures from extrasolar planets currently contributes to the growing interest about their nature and properties. Since the first planet discovered around another solar-type star in 1995 \\citep{Mayor:1995}, nearly 370 extrasolar planets have been detected and many more are expected to be unveiled by ongoing or future search programmes. Most extrasolar planets detected so far have been identified from the ground by indirect techniques, which rely on observable effects induced by the planet on its parent star. From the ground, radial velocity measurements are currently limited to the detection of planets about 2 times as massive as the Earth in orbits around Sun-like and low-mass stars \\citep{Mayor:2009b} while the transit method is limited to Neptune-sized planets \\citep{Gillon:2007}. Thanks to the very high precision photometry enabled by the stable space environment, the first space-based dedicated missions (namely CoRoT and Kepler) are now expected to reveal Earth-sized extrasolar planets by transit measurements as well. Launched in 2006, CoRoT (Convection Rotation and planetary Transits) has detected its first extrasolar planets \\citep[e.g.,][]{Barge:2008,Alonso:2008,Leger:2009} and is expected to unveil about 100 transiting planets down to a size of 2\\,R$_{\\oplus}$ around G0V stars and 1.1\\,R$_{\\oplus}$ around M0V stars over its entire lifetime for short orbital periods \\citep{Moutou:2005}. Launched in 2009, Kepler will extend the survey to Earth-sized planets located in the habitable zone of about 10$^5$ main sequence stars \\citep{Borucki:2007}. After 4 years, Kepler should have discovered several hundred of terrestrial planets with periods between one day and 400 days. After this initial reconnaissance by CoRot and Kepler, the Space Interferometry Mission (SIM PlanetQuest) might provide unambiguously the mass of Earth-sized extrasolar planets orbiting in the habitable zone of nearby stars by precise astrometric measurements. With CoRoT and Kepler, we will have a large census of Earth-sized extrasolar planets and their occurrence rate as a function of various stellar properties. However, even though the composition of the upper atmosphere of transiting extrasolar planets can be probed in favorable cases \\citep[e.g.,][]{Richardson:2007}, none of these missions will directly detect the photons emitted by the planets which are required to study the planet atmospheres and eventually reveal the signature of biological activity.   Detecting the light from an Earth-like extrasolar planet is very challenging due to the high contrast ($\\sim$10$^7$ in the mid-IR, $\\sim$10$^{10}$ in the visible) and the small angular separation ($\\sim$ 0.5 $\\mu$rad for an Earth-Sun system located at 10\\,pc) between the planet and its host star. A technique that has been proposed to overcome these difficulties is nulling interferometry \\citep{Bracewell:1978}. The basic principle is to combine the beams coming from two telescopes in phase opposition so that a dark fringe appears on the line of sight, which strongly reduces the stellar emission. Considering the two-telescope interferometer initially proposed by Bracewell, the response on the plane of the sky is a series of sinusoidal fringes, with angular spacing of $\\lambda/b$. By adjusting the baseline length ($b$) and orientation, the transmission of the off-axis planetary companion can then be maximised. However, even when the stellar emission is sufficiently reduced, it is generally not possible to detect Earth-like planets with a static array configuration, because their emission is dominated by the thermal contribution of warm dust in our solar system as well as around the target stars (exozodiacal cloud). This is the reason why Bracewell proposed to rotate the interferometer so that the planetary signal is modulated by alternatively crossing high and low transmission regions, while the stellar signal and the background emission remain constant. The planetary signal can then be retrieved by synchronous demodulation. However, a rapid rotation of the array would be difficult to implement and as a result the detection is highly vulnerable to low frequency drifts in the stray light, thermal emission, and detector gain. A number of interferometer configurations with more than two collectors have then been proposed to perform faster modulation and overcome this problem by using phase chopping \\citep{Angel:1997,Mennesson:1997,Absil:2001}.  The principle of phase chopping is to synthetize two different transmission maps with the same telescope array, by applying different phase shifts in the beam combination process. By differencing two different transmission maps, it is possible to isolate the planetary signal from the contributions of the star, local zodiacal cloud, exozodiacal cloud, stray light, thermal, or detector gain. Phase chopping can be implemented in various ways (e.g.\\ inherent and internal modulation, \\citealt{Absil:thesis}), and are now an essential part of future space-based life-finding nulling interferometry missions such as ESA's \\darwin{} \\citep{Fridlund:2006} and NASA's Terrestrial Planet Finder \\citep[TPF,][]{Lawson:2008}. The purpose of this paper is to assess the impact of exozodiacal dust discs on the performance of these missions. After describing the nominal performance of \\darwin/TPF, the first part of the study is dedicated to centrally symmetric exozodiacal discs which are suppressed by phase chopping and only contribute through their shot noise. If they are too bright, exozodiacal discs can drive the integration time and we investigate the corresponding impact on the number of planets that can be surveyed during the mission lifetime. In the second part, we address the impact of asymmetric structures in the discs (such as clumps and offset) which are not canceled by phase chopping and can seriously hamper the planet detection process. ",
        "conclusions": "Infrared nulling interferometry is the core technique of future life-finding space missions such as ESA's \\darwin{} and NASA's Terrestrial Planet Finder (TPF). Observing in the infrared (6-20 $\\mu$m), these missions will be able to characterise the atmosphere of habitable extrasolar planets orbiting around nearby main sequence stars. This ability to study distant planets strongly depends on exozodiacal clouds around the stars, which can hamper the planet detection. Considering the nominal mission architecture with 2-m aperture telescopes, we show that centrally symmetric exozodiacal dust discs about 100 times denser than the solar zodiacal cloud can be tolerated in order to survey at least 150 targets during the mission lifetime. The actual number of planet detections will then depend on the number of terrestrial planets in the habitable zone of target systems.  The presence of asymmetric structures in exozodiacal discs (e.g.,\\ clumps or offset) may be more problematic. While the cloud brightness drives the integration time necessary to disentangle the planetary photons from the background noise, the emission from inhomogeneities are not perfectly subtracted by phase chopping so that a part of the disc signal can mimic the planet. To address this issue, we consider modeled resonant structures produced by an Earth-like planet and introduce the corresponding image into \\darwinsim, the mission science simulator. Even for exozodiacal discs a few times brighter than the solar zodiacal cloud, the contribution of these asymmetric structures can be much larger than the planetary signal at the output of the interferometer. Fortunately, the high angular resolution provided by the long imaging baseline of \\darwin/TPF in the X-array configuration is sufficient to spatially distinguish most of the extended exozodi emission from the planetary signal and only the hole in the dust distribution near the planet significantly contributes to the noise level. Considering the full wavelength range of \\darwin/TPF, we show that the tolerable dust density is about 15 times the solar zodiacal density for face-on systems and decreases with the disc inclination. In practice, this constraint might be relaxed since we examined a resonant ring model that does not include dust from highly eccentric or inclined parent bodies, the effects of grain-grain collisions, or perturbations by additional planets, all of which can reduce the contrast of the resonant ring and improve the tolerance to the exozodiacal dust density.   These results show that asymmetric structures in exozodiacal discs around nearby main sequence stars are one of the main noise sources for future exo-Earth characterization missions. A first solution to get around this issue is to have a long imaging baseline architecture which resolves out the more spatially extended emission of the exozodiacal cloud from the point-like emission of planets. The stretched X-array configuration is particularly convenient in that respect. The second solution is to observe in advance the nearby main sequence stars and remove from the \\darwin/TPF target list those presenting a too high dust density or disc inclination. The FKSI nulling interferometer would be ideal in that respect with the possibility to detect exozodiacal discs down to the density of the solar zodiacal cloud. Ground-based nulling instruments like LBTI and ALADDIN would also be particularly valuable.  \\begin{figure}[!t] \\begin{center} \\includegraphics[width=8.5 cm]{zodi_vs_off.eps} \\caption{Tolerable exozodiacal dust density with respect to the offset between the center of symmetry of the exozodiacal disc and the central star (a G2V star located at 15\\,pc). The disc is assumed to be seen in face-on orientation.}\\label{Fig:zodi_vs_off} \\end{center} \\end{figure}  \\appendix"
    },
    "0910/0910.1356_arXiv.txt": {
        "abstract": "The existence of black holes of masses $\\sim 10^2$--$10^5 {\\rm M_{\\odot}}$ has important implications for the formation and evolution of star clusters and supermassive black holes. One of the strongest candidates to date is the hyperluminous X-ray source HLX1, possibly located in the S0-a galaxy ESO\\,243-49, but the lack of an identifiable optical counterpart had hampered its interpretation. Using the {\\it Magellan} telescope, we have discovered an unresolved optical source with $R = 23.80\\pm0.25$ mag and $V = 24.5\\pm0.3$ mag within HLX1's positional error circle. This implies an average X-ray/optical flux ratio $\\sim 500$. Taking the same distance as ESO\\,243-49, we obtain an intrinsic brightness $M_R = -11.0 \\pm 0.3$ mag, comparable to that of a massive globular cluster. Alternatively, the optical source is consistent with a main-sequence M star in the Galactic halo (for example an M4.4 star at $\\approx 2.5$ kpc). We also examined the properties of ESO\\,243-49 by combining {\\it Swift}/UVOT observations with stellar population modelling. We found that the overall emission is dominated by a $\\sim 5$ Gyr old stellar population, but the UV emission at $\\approx 2000$ \\AA\\ is mostly due to ongoing star-formation at a rate of $\\sim 0.03 {\\rm M_{\\odot}}$ yr$^{-1}$. The UV emission is more intense (at least a $9\\sigma$ enhancement above the mean) North East of the nucleus, in the same quadrant as HLX1. With the combined optical and X-ray measurements, we put constraints on the nature of HLX1. We rule out a foreground star and a background AGN. Two alternative scenarios are still viable. HLX1 could be an accreting intermediate mass black hole in a star cluster, which may itself be the stripped nucleus of a dwarf galaxy that passed through ESO\\,243-49, an event which might have caused the current episode of star formation. Or, it could be a neutron star in the Galactic halo, accreting from an M4--M5 donor star. ",
        "introduction": "{\\it XMM-Newton} and {\\it Chandra} have discovered several non-nuclear X-ray sources in nearby galaxies, with luminosities up to two orders of magnitude higher than those observed from Galactic X-ray binaries. These are referred to as ultraluminous X-ray sources \\citep[ULXs; e.g.,][]{ggs03,swa04,rob07}. Those findings have challenged our current models of black hole (BH) formation and accretion. Isotropic, Eddington-limited luminosities $\\ga 10^{40}$erg s$^{-1}$ would require BH masses $\\ga 100 {\\rm M_{\\odot}}$, beyond the upper limit for individual stellar collapses \\citep{yun08}. Mildly super-Eddington luminosity (possibly due to large super-Eddington mass accretion) from particularly heavy stellar BHs ($M \\sim 50 {\\rm M_{\\odot}}$), associated with mildly anisotropic emission may explain X-ray luminosities up to $\\sim$ few $\\times 10^{40}$erg s$^{-1}$ without the need for more exotic astrophysical processes \\citep{pou07,rob07,kin09}.  Only a few non-nuclear sources have been observed at X-ray luminosities $\\approx 0.7$--$1 \\times 10^{41}$erg s$^{-1}$. For example, in the Cartwheel \\citep{wol07}, in M\\,82 \\citep{fen09}, in NGC\\,2276 \\citep{dav04}, and in NGC\\,5775 \\citep{li08}. It is possible that such rare, extreme ULXs (sometimes known as hyperluminous X-ray sources, HLXs) may be powered by heavier BHs, formed through different channels: for example, in the collapsed core of a super star cluster, or within the nuclear star cluster of an accreted (and now disrupted) dwarf galaxy \\citep{kin05,bek03}. Thus, HLXs may represent evidence of intermediate-mass BHs. However, the debate is far from settled, given the small number of HLXs known, and the possibility of confusion with background AGN. The strongest claim for an X-ray luminous intermediate-mass BH so far has been made for a recently discovered X-ray source \\citep[2XMM J011028.1$-$460421, hereafter HLX1:][]{far09,god09} apparently located in the galaxy ESO\\,243-49, or, at least, projected inside the $\\mu_{B} = 25.0$ mag arcsec$^{-2}$ surface brightness isophote of that galaxy. Here, we report our discovery of the likely optical counterpart to this source, and our analysis of the UV emission in ESO\\,243-49. By determining the optical flux, and the X-ray/optical flux ratio, we test alternative models for the nature of this object. Our results strengthen the interpretation that the X-ray source belongs to ESO\\,243-49. We suggest that it is located inside a massive star cluster.  ESO\\,243-49 is an edge-on S0-a galaxy at a luminosity distance of $91 \\pm 6$ Mpc \\citep[$z=0.0224$, distance modulus $34.80 \\pm 0.15$:][]{cal97}. The foreground extinction is very low, $A_V = 0.043$ mag \\citep{sch98}. HLX1 appears projected $\\approx 7\\arcsec$ ($\\approx 3.1$ kpc) to the North-East of ESO\\,243-49's nucleus, and $\\approx 1\\farcs8$ ($\\approx 800$ pc) above the galactic plane. HLX1 has been detected several times with {\\it XMM-Newton}, {\\it Chandra} and {\\it Swift} between 2004 and 2009 \\citep{far09,god09}, with an unabsorbed luminosity in the $0.3$--$10$ keV band varying between $\\la 5 \\times 10^{40}$ erg s$^{-1}$ and $\\approx 1 \\times 10^{42}$ erg s$^{-1}$. We also examined a {\\it ROSAT}/HRI observation of the field from 1996, when HLX1 was not detected to an upper limit of $\\approx 5 \\times 10^{40}$ erg s$^{-1}$. Its combination of extreme luminosity (if it really belongs to ESO\\,243-49), soft spectrum, and spectral changes on short timescales \\citep{god09} makes it a unique object among the ULX/HLX class. Its apparent location in an S0 galaxy is also puzzling, because such galaxies are usually dominated by an old stellar population. For example, the integrated brightnesses of ESO\\,243-49 are (Cousins) $B=14.92 \\pm 0.09$ mag, (Cousins) $R=13.48 \\pm0.09$ mag and (2MASS) $K = 10.70 \\pm 0.05$ mag (from NED\\footnote{http://nedwww.ipac.caltech.edu}), which are indicative of a characteristic stellar age $\\sim$ a few Gyr (Section 4). Such moderately old populations were not previously known to host luminous ULXs or HLXs. For these reasons, it was speculated that the source might be a background AGN or a foreground neutron star, even though its X-ray properties are also very unusual for both classes of objects (Section 5).     \\begin{figure} \\psfig{figure=f1_xv.ps,width=84mm} \\caption{{\\it Magellan}/IMACS $R$-band image of ESO\\,243-49, with the X-ray position of HLX1 marked by a circle of $0\\farcs5$ radius (combined astrometric uncertainty of the X-ray and optical images).} \\label{f1} \\end{figure}    \\begin{table} \\begin{center} \\begin{tabular}{llrr} \\hline Telescope & Date  & Band & Exposure \\\\ \\hline {\\it Magellan} Baade & 2009 Aug 26 & R & 540 s  \\\\ &                         & V & 540 s \\\\ \\hline {\\it Swift}/UVOT & 2008 Oct 24 & $u$ &  380 s\\\\ &                         & $uvw1$ & 760 s \\\\ &                         & $uvw2$ & 196 s \\\\ & 2008 Oct 25 & $uvw2$ &   1264 s\\\\ & 2008 Nov 01 & $u$ &   730 s\\\\ &                         & $uvw1$ & 1690 s \\\\ &                         & $uvw2$ & 2639 s \\\\ & 2008 Nov 07  & $u$ &   1210 s\\\\ &                         & $uvw1$ & 2410 s \\\\ &                         & $uvw2$ & 3278 s \\\\ & 2008 Nov 08 & $uvw2$ &   582 s\\\\ & 2008 Nov 14 & $u$ &   980 s\\\\ &                         & $uvw1$ & 1960 s \\\\ &                         & $uvw2$ & 3814 s \\\\ & 2009 Aug 05 & $uvw2$ &  9753 s\\\\ & 2009 Aug 06 & $uvw2$ &   9132 s\\\\ & 2009 Aug 16 & $uvw2$ &  5664 s\\\\ & 2009 Aug 17 & $uvw2$ &   681 s\\\\ & 2009 Aug 18 & $uvw2$ &   6032 s\\\\ & 2009  Aug 19 & $uvw2$ &   4257 s\\\\ & 2009 Aug 20 & $uvw2$ &   2199 s\\\\ & 2009  Nov 02 & $uvw2$ & 8956 s\\\\ & 2009  Nov 14 & $uvw2$ & 3903 s\\\\ & 2009  Nov 20 & $uvw2$ & 517 s\\\\ & 2009  Nov 21 & $uvw2$ & 1453 s\\\\ & 2009  Nov 28 & $uvw2$ & 2981 s\\\\ & 2009  Nov 29 & $uvw2$ & 2028 s\\\\ & 2009  Dec 05 & $uvw2$ & 2993   s\\\\ & 2009  Dec 19 & $uvw2$ & 2518   s\\\\ & 2009  Dec 26 & $uvw2$ & 2774    s\\\\ & 2010  Jan 02 & $uvw2$ & 2590   s\\\\ & 2010  Jan 08 & $uvw2$ & 4348   s\\\\ & 2010  Jan 13 & $uvw2$ & 3577    s\\\\ & 2010  Jan 15 & $uvw2$ & 3068    s\\\\ & 2010  Jan 22 & $uvw2$ & 2945   s\\\\ \\hline \\end{tabular} \\end{center} \\caption{Optical/UV observation log. } \\end{table}        \\begin{figure*} \\psfig{figure=new_f2a_xv.ps,width=87mm} \\psfig{figure=new_f2b.ps,width=87mm}\\\\ \\psfig{figure=new_f2c.ps,width=87mm} \\psfig{figure=new_f2d.ps,width=87mm} \\caption{{\\it Top left:} differential $R$-band image of ESO\\,243-49, with a median-filter smoothed image subtracted from the original image (logarithmic greyscale). The source marked with an arrow has a $4\\sigma$ significance and is located near the X-ray position of HLX1. {\\it Top right:} zoomed-in view of the field around HLX1, in the $R$-band residual image (square-root greyscale). The {\\it Chandra}/HRC-I position of HLX1 is marked by a circle of $0\\farcs5$ radius (combined astrometric uncertainty of the X-ray and optical images). Both the top and bottom panel images have been Gaussian-smoothed with a kernel radius of 2 pixels ($0\\farcs22$), for display purposes only. {\\it Bottom left:} Gaussian-smoothed field around HLX1, in the $V$-band residual image (logarithmic greyscale); an optical counterpart is located at the same position with a $3\\sigma$ significance. {\\it Bottom right:} Gaussian-smoothed field around HLX1, from the combined $R$-band plus $V$-band residual images.} \\label{f2} \\end{figure*} ",
        "conclusions": "If the X-ray source HLX1 is proven to be an accreting BH with mass $\\sim 10^3$--$10^4 {\\rm M_{\\odot}}$, there would be important implications on models of galaxy formation and evolution. Identifying its optical counterpart gives a crucial constraint on its nature. We have found an unresolved optical source within its X-ray error circle, and it is likely to be physically associated to HLX1. Assuming a direct association, we calculate an X-ray/optical flux ratio, using the standard definition \\citep{hor01}: $\\log(f_{\\rm X}/f_{\\rm R}) = \\log f_{\\rm X} + 5.5 + R/2.5$, where $f_{\\rm X}$ is the intrinsic flux in the $0.3$--$10$ keV band \\citep[taken from][and our spectral analysis of the {\\it Swift}/XRT data]{god09}. We obtain $f_{\\rm X}/f_{\\rm R} \\approx 800$--$1000$ for the X-ray high state of 2009 August, and $f_{\\rm X}/f_{\\rm R} \\approx 500$ for the ``average'' X-ray state where the source was more often observed over 2008--2009; such ratio is only slightly dependent on the choice of X-ray spectral model. The observed X-ray/optical flux ratio is much higher than expected for AGN, which have typical $f_{\\rm X}/f_{\\rm R} \\la 10$ \\citep{lai09,bau04}. A number of distant, faint AGN are undetected in the optical band because of extinction, which is not an issue for this object \\citep{far09,god09}. In particular, AGN with a $0.5$--$2$ keV flux of $\\approx 5 \\times 10^{-13}$ erg s$^{-1}$ cm$^{-2}$ are rare \\citep[$N \\approx 0.5$ deg$^{-2}$:][]{has98} and have always been easily identified in other bands. Moreover, the red colour of the optical counterpart ($V-R = 0.7 \\pm 0.4$) suggests that the optical emission is not dominated by the Rayleigh-Jeans tail of an accretion disk spectrum.  Neutron stars are a class of objects that can reach X-ray/optical flux ratios $\\ga 1000$, with thermal X-ray emission from their surface. We note (Soria et al., in prep.) that the X-ray spectra of HLX1 can also be fitted with a neutron star atmosphere model (e.g., {\\tt nsa} in {\\footnotesize XSPEC}) plus power-law. Such models suggest a characteristic distance $\\approx 1.5$--$2.5$ kpc, placing the source in the Galactic halo. The corresponding $0.3$--$10$ keV luminosity would be $\\approx$ a few $\\times 10^{32}$ erg s$^{-1}$. Intriguingly, this is a range of luminosities where quiescent low-mass X-ray binaries also show a thermal plus power-law spectrum, with the relative contribution of the power-law decreasing as the source gets brighter \\citep{jon04}. In this scenario, the apparent optical brightness $R \\approx 23.8$ mag implies $M_R \\approx 11.8$--$12.8$ mag, consistent with a main-sequence M4.4--M5.2 donor star, with an initial mass $\\approx 0.13$--$0.17 {\\rm M_{\\odot}}$ \\citep{kni06,gir00}. A late-type M star is also consistent with the observed red colour. Thus, we cannot rule out the possibility of an accreting neutron star in the Galactic halo, from the optical data. In this case, the excess UV emission North East of the nucleus in ESO\\,243-49 is purely coincidental.   The optical counterpart of HLX1 does not stick out like a unique object in the field. In the $R$-band residual image, we identified a few other sources consistent with globular clusters around ESO\\,243-49, with comparable brightnesses (Figure 2, top panel). If the HLX1 counterpart is also a globular cluster in that galaxy, its optical luminosity would place it between the Milky Way globular cluster $\\omega$ Cen \\citep[$M_V = -10.3$ mag, $M_R = -10.8$ mag, $M_{\\rm tot} \\approx 2.8 \\times 10^6 {\\rm M_{\\odot}}$:][]{har96,van09} and the Andromeda cluster G1 \\citep[$M_V = -11.2$ mag, $M_R = -11.8$ mag, $M_{\\rm tot} \\approx 5 \\times 10^6 {\\rm M_{\\odot}}$:][]{gra09}, which is a strong candidate for the presence of an $\\approx 2 \\times 10^4 {\\rm M_{\\odot}}$ BH \\citep[][but see \\citealt{bau03}]{geb05}.  There are a number of scenarios for the formation of an intermediate-mass BH inside a massive star cluster. In a young star cluster, runaway core collapse and coalescence of the most massive stars can occur over a timescale $\\la 3$ Myrs and can result in the formation of a supermassive star, which can quickly collapse into a BH \\citep{por02,fre06}. In an old globular cluster, an intermediate-mass BH can be formed from the merger of stellar-mass BHs and neutron stars over a timescale $\\sim 10^9$ yr \\citep{ole06}. Subsequent capture and disruption of a cluster star may then provide the accretion rate required to explain the X-ray luminosity. In this scenario, HLX1 and its host star cluster are unrelated to the ongoing star formation in the bulge of ESO\\,243-49.  The most intriguing scenario is that some massive star clusters may have been the nuclear clusters of satellite galaxies accreted and tidally disrupted by a more massive galaxy. Dwarf galaxies are the most common type of galaxies in clusters \\citep[e.g.,][]{bin85} and many of them are nucleated \\citep[e.g.,][]{gra03,cot06}. In many cases, a nuclear cluster may coexist with a nuclear BH \\citep{gra09,set08}. This may end up in the halo of a bigger galaxy after a merger. $\\omega$ Cen itself may have originated from the nuclear star cluster of an accreted dwarf \\citep{bek03}. Similar suggestions have been made for a group of clusters in NGC\\,5128 \\citep{pen02,cha09}. The recent or ongoing star formation in ESO\\,243-49 may have been triggered by the passage and tidal disruption of the satellite galaxy, perhaps along the South-West to North-East direction, since {\\it uvw2} emission is stronger on that side (Section 4). During its passage through ESO\\,243-49, the compact nucleus of the satellite dwarf may also have collected gas from the main galaxy \\citep{pfl09}, and this may perhaps be fuelling a nuclear BH. In this scenario, HLX1 may be the intermediate-mass BH located in the nuclear cluster of that accreted satellite.  In summary, we have identified a point-like optical counterpart for HLX1 in ESO\\,243-49 in the {\\it Magellan} images. The optical brightness and colour are consistent with a massive star cluster in ESO\\,243-49, or with a main-sequence M4--M5 star in the Milky Way halo, at $\\approx 1.5$--$2.5$ kpc. The galaxy is dominated by a $\\sim 5$ Gyr old population, but shows excess emission in the {\\it Swift}/UVOT $uvw2$ band, consistent with a recent episode of star formation. The far-UV emission has an asymmetric shape and is stronger to the North East of the nucleus, roughly in the direction of HLX1."
    },
    "0910/0910.4745_arXiv.txt": {
        "abstract": "Particle cascades initiated by ultra-high energy (UHE) neutrinos in the lunar regolith will emit an electromagnetic pulse with a time duration of the order of nano seconds through a process known as the Askaryan effect. It has been shown that in an observing window around 150~MHz there is a maximum chance for detecting this radiation with radio telescopes commonly used in astronomy. In 50 hours of observation time with the Westerbork Synthesis Radio Telescope array we have set a new limit on the flux of neutrinos, summed over all flavors, with energies in excess of $4\\times10^{22}$~eV. ",
        "introduction": "At high energies the spectrum of cosmic rays follows a power law distribution extending up to extremely large energies. At the Pierre Auger Observatory cosmic rays have been observed~\\cite{A08} with energies in excess of $\\sim 10^{20}$~eV. Above the Greisen-Zatsepin-Kuzmin (GZK) energy of $6\\times 10^{19}$~eV, cosmic rays can interact with the photons of the cosmic microwave background to produce pions~\\cite{G66,ZK66} which carry a sizable fraction of the original energy of the cosmic ray. Charged pions decay and produce neutrinos and one thus may expect the presence of neutrinos with energies in excess of the GZK energy. Recently at the Pierre Auger Observatory a steepening of the slope in the cosmic ray spectrum has been observed at the GZK energy~\\cite{A08} which can be regarded as a consequence of the GZK effect.  Since neutrinos are chargeless they will propagate in a straight line with negligible energy loss from the location where they have been created to the observer, thus carrying direct information on their source. These sources could be the aforementioned processes related to the GZK effect or, more exotically, decaying supermassive dark-matter particles or topological defects. This last class of models is referred to as top-down (see \\citet{s04} for a review).  Because of their small interaction cross section, the detection of cosmic neutrinos calls for extremely large detectors. At GeV energies their cross section is so minute that at a flux given by the Waxman-Bahcall estimate~\\cite{Waxman:1998yy} of a few tens of neutrinos per km$^2$ per year one needs to employ km$^3$-scale detectors~\\cite{icecube,KM3NET}. At higher neutrino energies the reaction cross section increases. However, their flux is expected to fall even faster and one needs even larger detection volumes. These can be obtained by observing large detector masses from a distance. The ANITA balloon mission~\\cite{anita08} monitors an area of a million km$^2$ of South Pole ice and the FORTE satellite~\\cite{forte} can pick up radio signals coming from the Greenland ice mass. Alternatively the Pierre Auger Observatory can distinguish cosmic ray induced air showers from neutrino induced cascades at very high zenith angles~\\cite{Abraham:2007rj}.  The Moon offers an even larger natural detector volume. In the interaction of an UHE neutrino about 20\\% of the energy of the neutrino is converted into a cascade of high-energy particles, called the hadronic shower. Due to the electromagnetic component of this shower the electrons in the material are swept out from the atom to become part of the shower. The shower thus has an excess of negative charge moving with relativistic velocities through a material with an index of refraction considerably different from unity resulting in the emission of Cherenkov radiation. Since the lateral side of the shower has a dimension of the order of 10~cm, the radio emission is coherent for wavelength of this magnitude and larger or for frequencies up to $\\sim 3$ GHz. The emission of coherent Cherenkov radiation in such a process is known as the Askaryan effect~\\cite{a62}. This emission mechanism has been experimentally verified at accelerators~\\cite{s01} and extensive calculations have been performed to quantify the effect~\\cite{zhs92,az97}.  For showers in the lunar regolith, the top layer of the Moon consisting of dust and small rocks, its properties are important. Much is known from samples brought from the Moon~\\cite{os75}. The average index of refraction is $n=1.8$ and the attenuation length is $\\lambda_{r}=(9/\\nu[\\mathrm{GHz}])$~m for radio waves. The thickness of the regolith is known to vary over the lunar surface. At some depth there is a (probably smooth) transition to solid rock, for which the density is about twice that of the regolith.  At the highest energies the emitted pulse from the Moon can be observed at Earth with radio telescopes~\\cite{dz89}. The first experiments in this direction were carried out with the Parkes telescope~\\cite{parkes}, later followed by others~\\cite{glue,KALYAZIN}. A recent project is LUNASKA~\\cite{James:2009rc} that is currently performing lunar Cherenkov measurements with the Australia Telescope Compact Array with a 600 MHz bandwidth at 1.2-1.8~GHz. These observations are all performed at relatively high frequencies (2~GHz) where the emission is strongest. At lower frequencies~\\cite{Fal03} the angular spread of the emission around the Cherenkov angle increases due to finite source effects~\\cite{scholten}. When the wavelength is similar to the longitudinal extent of the shower in the lunar rock, a few meters, the angular spread is close to isotropic and the probability of detecting the radio pulse is largest~\\cite{scholten}. In our observations we exploit this optimal frequency range around 150 MHz using the Westerbork Synthesis Radio Telescope (WSRT). ",
        "conclusions": ""
    },
    "0910/0910.3023_arXiv.txt": {
        "abstract": "We present initial results from a new 440-ks Chandra HETG GTO observation of the canonical Seyfert 2 galaxy NGC 1068. The proximity of NGC 1068, together with Chandra's superb spatial and spectral resolution, allow an unprecedented view of its nucleus and circumnuclear NLR. We perform the first spatially resolved high-resolution X-ray spectroscopy of the `ionization cone' in any AGN, and use the sensitive line diagnostics offered by the HETG to measure the ionization state, density, and temperature at discrete points along the ionized NLR. We argue that the NLR takes the form of outflowing photoionized gas, rather than gas that has been collisionally ionized by the small-scale radio jet in NGC 1068. We investigate evidence for any velocity gradients in the outflow, and describe our next steps in modeling the spatially resolved spectra as a function of distance from the nucleus. ",
        "introduction": "Outflows and feedback from AGN are widely invoked as the key mediator between the co-evolution of black holes and their host galaxies over cosmic time. It is now well understood from large optical surveys that galaxies evolve through mergers from blue, star-forming spirals (the `blue cloud') whose black holes accrete at or near their Eddington limits, through a transition region (the `green valley'), to so-called `red and dead' ellipticals (the `red sequence'), which instead are described by markedly less black-hole growth. The importance of outflows in this evolution has now taken center stage; yet before we can successfully incorporate AGN feedback into numerical simulations of galaxy growth, a number of key questions need to be answered from observations:  \\begin{itemize} \\item Can AGN actually deliver enough power to their environments to alter the evolution of their host galaxies in a meaningful way? \\item On what spatial scales does this occur? \\item Where are the outflows in the first place? \\end{itemize}  Several steps have recently been taken to address these issues. It is now apparent from multiwavelength surveys that AGN selected in different wavelengths tend to populate different regions of the galaxy color-magnitude diagram. For example, \\cite{hic09} showed that IR-selected AGN are characterized by high Eddington ratios and tend to inhabit the `blue cloud'; X-ray AGN have intermediate-to-high ratios and lie in the `green valley'; and radio AGN have low Eddington fractions and populate the `red sequence'. The relative paucity of galaxies in the `green valley', together with the prevalence of X-ray AGN in this region suggest that the kpc-scale, ionized Narrow-Line Regions (NLR) in such AGN are ideal places to search for outflows, together with their kinematic and ionizing effect on their galaxy-scale environments. ",
        "conclusions": ""
    },
    "0910/0910.0216_arXiv.txt": {
        "abstract": "We investigate the solar flare of 20 October 2002. The flare was accompanied by quasi-periodic pulsations (QPP) of both thermal and nonthermal hard X-ray emissions (HXR) observed by RHESSI in the 3-50 keV energy range. Analysis of the HXR time profiles in different energy channels made with the Lomb periodogram indicates two statistically significant time periods of about 16 and 36 seconds. The 36-second QPP were observed only in the nonthermal HXR emission in the impulsive phase of the flare. The 16-second QPP were more pronounced in the thermal HXR emission and were observed both in the impulsive and in the decay phases of the flare. Imaging analysis of the flare region, the determined time periods of the QPP and the estimated physical parameters of magnetic loops in the flare region allow us to interpret the observations as follows. \\textbf{1)} In the impulsive phase energy was released and electrons were accelerated by successive acts with the average time period of about 36 seconds in different parts of two spatially separated, but interacting loop systems of the flare region. \\textbf{2)} The 36-second periodicity of energy release could be caused by the action of fast MHD oscillations in the loops connecting these flaring sites. \\textbf{3)} During the first explosive acts of energy release the MHD oscillations (most probably the sausage mode) with time period of 16 seconds were excited in one system of the flare loops. \\textbf{4)} These oscillations were maintained by the subsequent explosive acts of energy release in the impulsive phase and were completely damped in the decay phase of the flare. ",
        "introduction": "\\label{S-Introduction}  \\indent Electrons, accelerated during solar flares, generate nonthermal HXR and microwave emission in the solar atmosphere as a result of their Coulomb collisions with ambient plasma (bremsstrahlung) and interaction with magnetic field (synchrotron radiation), respectively. Often the light curves of nonthermal HXR and microwave emissions in large solar flares are composed of multiple pulses with different durations from less than one second up to several minutes (\\eg,~ \\opencite{dennis88};~\\opencite{aschwanden09}, and references therein). This indicates many episodes of acceleration of electrons, possibly in different locations within flare regions, since the majority of these flares are accompanied by formation of arcades of many flaring loops and by the HXR emission generated in different flare loops (\\eg,~ \\opencite{vorpahl76};~\\opencite{dejager79};~\\opencite{grigis05}). It has been suggested that flares consist of a number of ``Elementary Flare Bursts'' (EFB), which can be a result of collisions between different current-carrying loops \\cite{dejager79,sakai96}.  \\indent More rarely, the light curves of nonthermal HXR and microwave emission show apparent quasi-periodic pulsations (QPP; see, \\eg,~ \\opencite{dennis88}; ~\\opencite{nakariakov05};~ \\opencite{aschwanden09} for a review and references therein). This may indicate the presence of some quasi-periodic processes (not fully understood at the moment) of acceleration of electrons. Using RHESSI imaging observations in the HXR range, it was shown that this quasi-periodic acceleration of electrons, at least in some flares, occurs in different flare loops, rather than in a single one \\cite{zimovets09}. In this case the observed quasi-periodicity could be a coincidence owing to similar physical conditions in successively bursting flare loops. This returns us to the idea of multiple EFB due to coalescence of current-carrying loops or due to the presence of multiple spatially separated ``null-points'' in flare regions, in which the stored magnetic energy can be released. But in principle, it is possible to implement a model proposed by \\inlinecite{nakariakov06} to explain the QPP in the latter case by the successive quasi-periodic reconnection in different null-points due to fast magnetoacoustic oscillations in nearby non-flaring loops.  \\indent On the other hand, it was theoretically suggested that QPP could be produced inside a single vibrating flare loop by the quasi-periodic modulation of the spectrum of nonthermal electrons through the betatron acceleration \\cite{brown75} or by the quasi-periodic modulation of efficiency of the trapped nonthermal electrons to precipitate into the lower heights with larger plasma density and stronger magnetic field (in particular, the sausage mode is a good candidate for this) \\cite{zaitsev82}. Indeed, imaging observations of the QPP in the microwave range made using the Nobeyama Radioheliograph with sufficient spatial resolution (but with no simultaneous QPP observed in the HXR range), which could be interpreted in terms of the global sausage mode of the flaring loop oscillations, were reported \\cite{nakariakov03}. Imaging observations of QPP, but without sufficient spatial resolution together with an analysis of the time periods of QPP, allowed QPP to be interpreted in terms of the flare loop oscillations (\\eg,~ \\opencite{asai01};~ \\opencite{inglis09}). Different modes of oscillations in coronal loops actually have been found with the TRACE and SOHO spacecraft in the range of the extreme ultraviolet radiation (see, \\eg,~ \\opencite{aschwanden09};~ \\opencite{nakariakov05} for a review and references therein), making the oscillations of magnetic loops a natural way to explain the QPP of electromagnetic emission in flares. However, despite the large number of papers published on flare QPP, their nature is still not fully understood. Further imaging analysis of flares with QPP is required.  \\indent In this work we investigate the solar flare of 20 October 2002. The light curves of its thermal and nonthermal HXR emission detected by RHESSI clearly indicate the presence of the QPP with two significant time periods of about 16 and 36 seconds. We will present time-spectral, energy-spectral and imaging analysis of the flare HXR emission and propose possible flare scenario to interpret the observed QPP. ",
        "conclusions": "% \\label{S-discussion}  Based on the analyzed observational data, we can propose a rough scenario of the investigated flare [see Figure~\\ref{fig3}(p)].  \\indent \\textbf{1.} \\textbf{In the impulsive phase of the flare energy was released and electrons were accelerated by successive quasi-periodic acts with an average time period of about 36 seconds in different parts of two spatially separated, but interacting loop-systems of the flare region.} Possibly, these loop-systems interact through the connecting magnetic loop near both ends of which they are located. We have established two observational evidences that these flare systems are linked: 1) the fluxes of thermal HXR emission from these systems are strongly correlated; 2) the formation of the loop-type structure in the space between these flaring systems is observed in thermal HXR emission. The question appears: is it possible to interpret the observed 36-second quasi-periodicity ($P_{36}$) of energy release in two different flaring systems by the MHD oscillations in the connecting loop? Indeed, it was concluded from the imaging observations of another flare event that oscillations of even nonflaring transequatorial loop (the kink modes) can be responsible for the QPP of HXR emission observed simultaneously from two different active regions located at both ends of this loop \\cite{foullon05}. Unfortunately, it is not possible to estimate length of the connecting loop ($L_{c}$) accurately in our case, due to insufficient spatial resolution of observations, nevertheless we can say with certainty that $20\\leq{L_{c}}\\leq40$ Mm. Thus, we can estimate the required phase speed of a fundamental standing mode as $V_{ph}=2L_{c}/P_{36}$ or $1100\\leq{V_{ph}}\\leq{2200}$ km/s. These phase speeds are consistent with the speeds of fast magnetoacoustic kink modes that have been directly observed in coronal loops in the EUV emission (\\opencite{nakariakov05}; \\opencite{aschwanden09}; and references therein). Implementing the imaging spectroscopy technique to the RHESSI HXR data we could estimate the emission measure of different parts of the flare region. Hence, using the observed volumes of these parts, we could roughly estimate averaged electron plasma density ($n_{e}$) in them. Our estimations give ${10^{10}}\\leq{n_{e}}\\leq{10^{11}}$ cm$^{-3}$ for each flaring loop-system and also for the connecting loop-type structure. Thus, the Alfv\\'{e}n speed in the connecting loop-type structure is $7B\\leq{V_{A}}\\leq{22B}$ ~km/s, where $B$ ~is magnetic field in gauss inside the loop-type structure. Consequently, if $B\\approx{100}$ ~G, which is a reasonable value, we have $700\\leq{V_{A}}\\leq{2200}$ ~km/s. These values of the Alfv\\'{e}n speed are very close to the estimated $V_{ph}$! Thus, we conclude that the observed 36-second periodicity can actually be caused by the fast MHD oscillations in the loop-type structure, which connects two separate flare systems. The physical mechanism, proposed by \\inlinecite{nakariakov06}, is a good candidate to explain quasi-periodic spatial fragmentation of energy release in the impulsive phase of this flare (see Section~\\ref{S-Introduction}). Multiple null-points could actually exist in this flare region, because of its complex topology: several inclusions of magnetic field of opposite polarity are seen on the magnetograms.  \\indent \\textbf{2.}\t\\textbf{During the first 36-second explosive episodes of energy release, the MHD oscillations, probably global sausage modes} (see below)\\textbf{, with time period of about 16 seconds were excited in the loops of the NW flare system.}  \\indent \\textbf{3.}\t\\textbf{These oscillations were maintained by the subsequent explosive acts of energy release in the impulsive phase and were damped in the decay phase of the flare.} We conclude that the 16-second oscillations were excited only in the NW system, but not in the SE flare system and not in the connecting loop-type structure at least in the decay phase, because the NW system was more stationary and retained its loop-like configuration in the decay phase, when the QPP of thermal HXR emission were still observed, but the SE flare system and the connecting loop-type structure were no longer visible in thermal HXR emission on the RHESSI images. Estimated length and electron plasma density of the loops in the NW flare system are $25\\leq{L_{NW}}\\leq{35}$ Mm and ${10^{10}}\\leq{n_{e}}\\leq{10^{11}}$~cm$^{-3}$, respectively. These physical parameters highly coincide with the same parameters, found in \\inlinecite{nakariakov03}, where the 16-second quasi-periodic pulsations of microwave emission, observed with the Nobeyama Radioheliograph, were successfully interpreted in terms of the global sausage mode of the oscillating flare loop! Thus, by analogy, we can also interpret the observed 16-second pulsations of HXR emission as the global sausage mode, excited in the loops of the NW flare system. Moreover, the sausage mode must be accompanied by perturbations of plasma density and emission measure. Hence, this mode of oscillations would quasi-periodically modulate flux of thermal HXR emission. This is what we observe both in the impulsive and in the decay phase of the flare. The decay phase was not accompanied by significant energy release, therefore we clearly observe the damping of oscillations. Contrary to that, in the impulsive phase there were multiple explosive acts of energy release, which might intermittently re-excite oscillations.  \\indent It is interesting to note that the 16-second periodicity of nonthermal emission has already been observed in other solar flares \\cite{parks69,nakariakov03,inglis08}. Perhaps, this is a typical period of the global sausage oscillations in flare regions \\cite{inglis08}. The 36-second periodicity of HXR emission was also found in many solar flares \\cite{lipa78}.  \\indent Despite the above arguments we are aware that the proposed flare scenario, based on the MHD oscillations of magnetic loops in flare region, may not be true. Oscillations of the flaring loops itself were not detected directly using imaging observations, though this may be due to the observational limitations.   \\begin{acks} We are grateful to the spacecraft teams and consortia (RHESSI, \\\\ SOHO, and GOES) and ground-based observatories (the Sagamore Hill Radio Observatory and the San Vito Solar Observatory), whose data were used in this work. We thank Jenny Harris for help in English editing. IVZ is grateful to the organizers of the Baikal Young Scientists' International School for Fundamental Physics (2009), where this work was presented. This work was partially supported by the Russian Foundation for Basic Research, grants No. 07-02-00319, 09-02-16032.  \\indent This article is dedicated to IVZ's grandfather F.I.~Krylov, veteran of the WWII, who died during its writing.  \\end{acks} \\clearpage{}"
    },
    "0910/0910.2025_arXiv.txt": {
        "abstract": "{} {Giant radio halos are mega-parsec scale synchrotron sources detected in a fraction of massive and merging galaxy clusters. Radio halos provide one of the most important pieces of evidence for non-thermal components in large scale structure. Statistics of their properties can be used to discriminate among various models for their origin. Therefore, theoretical predictions of the occurrence of radio halos are important as several new radio telescopes are about to begin to survey the sky at low frequencies with unprecedented sensitivity.} {In this paper we carry out Monte Carlo simulations to model the formation and evolution of radio halos in a cosmological framework. We extend previous works on the statistical properties of radio halos in the context of the turbulent re-acceleration model.} {First we compute the fraction of galaxy clusters that show radio halos and derive the luminosity function of radio halos. Then, we derive differential and integrated number count distributions of radio halos at low radio frequencies with the main goal to explore the potential of the upcoming LOFAR surveys. By restricting to the case of clusters at redshifts $<$ 0.6, we find that the planned LOFAR all sky survey at 120 MHz is expected to detect about 350 giant radio halos. About half of these halos have spectral indices larger than 1.9 and substantially brighten at lower frequencies. If detected they will allow for a confirmation that turbulence accelerates the emitting particles. We expect that also commissioning surveys, such as MS$^3$, have the potential to detect about 60 radio halos in clusters of the ROSAT Brightest Cluster Sample and its extension (eBCS). These surveys will allow us to constrain how the rate of formation of radio halos in these clusters depends on cluster mass.} {} ",
        "introduction": "Radio halos are diffuse Mpc--scale radio sources observed at the center of $\\sim 30\\%$ of massive galaxy clusters (\\egg Feretti 2005; Ferrari et al. 2008, for reviews). These sources emit synchrotron radiation due to GeV electrons diffusing through $\\mu$G magnetic fields and provide the most important evidence of non-thermal components in the intra-cluster-medium (ICM).  Clusters hosting radio halos always show very recent or ongoing merger events (\\eg Buote 2001; Schuecker et al 2001; Govoni et al. 2004; Venturi et al. 2008). This suggests a connection between the gravitational process of cluster formation and the origin of these halos. Cluster mergers are expected to be the most important sources of non-thermal components in galaxy clusters. A fraction of the energy dissipated during these mergers could be channelled into amplification of the magnetic fields (\\eg Dolag et al. 2002; Br\\\"uggen et al. 2005; Subramanian et al. 2006; Ryu et al. 2008) and into the acceleration of high energy particles via shocks and turbulence (\\eg En\\ss lin et al. 1998; Sarazin 1999; Blasi 2001; Brunetti et al. 2001, 2004; Petrosian 2001; Miniati et al. 2001; Fujita et al. 2003; Ryu et al. 2003; Hoeft \\& Br\\\"uggen 2007; Brunetti \\& Lazarian 2007; Pfrommer et al. 2008, Brunetti et al. 2009).   A promising scenario proposed to explain the origin of the synchrotron emitting electrons in radio halos assumes that electrons are re-accelerated due to the interaction with MHD turbulence injected in the ICM in connection with cluster mergers ({\\it turbulent re-acceleration} model, \\eg Brunetti et al. 2001; Petrosian 2001). An alternative possibility is that the emitting electrons are continuously injected by {\\it pp} collisions in the ICM ({\\it secondary} models; \\eg Dennison 1980; Blasi \\& Colafrancesco 1999).  In the picture of the {\\it turbulent re-acceleration} scenario, the formation and evolution of radio halos are tightly connected with the dynamics and evolution of the hosting clusters. Indeed, the occurrence of radio halos at any redshift depends on the rate of cluster-cluster mergers and on the fraction of the merger energy channelled into MHD turbulence and re-acceleration of high energy particles. In the last few years this has been modeled through Monte Carlo procedures (Cassano \\& Brunetti 2005; Cassano et al. 2006a) that provide predictions that can be verified using future instruments. In this scenario radio halos have a relatively short lifetime ($\\approx$ 1 Gyr), and the fraction of galaxy clusters where radio halos are generated is expected to increase with cluster mass (or X-ray luminosity), since the energy of the turbulence generated during cluster mergers is expected to scale with the cluster thermal energy (which roughly scales as $\\sim M^{5/3}$; \\eg Cassano \\& Brunetti 2005). It has been shown that the predicted occurrence of radio halos as a function of the cluster mass (or X-ray luminosity) is in line with results obtained from a large observational project, the ``GMRT radio halo survey'' (Venturi et al.~2007, 2008), and with its combination with studies of nearby halos based on the NVSS survey (\\eg Cassano et al.~2008).  \\begin{figure} \\centerline{ \\includegraphics[width=7.3cm,height=6.5cm]{USSRH_paper2.ps} } \\caption{A schematic representation of the theoretical synchrotron spectra of radio halos with different values of $\\nu_s$ (and $\\nu_b$). The colored regions indicate the frequency range of VLA and LOFAR observations.} \\label{fig:ussrh} \\end{figure}  \\noindent The steep spectrum of radio halos makes these sources ideal targets for observations at low radio frequencies suggesting that present radio telescopes can only detect the tip of the iceberg of their population (En\\ss lin \\& R\\\"ottgering 2002; Cassano et al.~2006; Hoeft et al. 2008). The recent discovery of the giant and ultra-steep spectrum radio halo in Abell 521 at low radio frequencies (Brunetti et al. 2008) allows a first confirmation of this conjecture and provides a glimpse of what future low frequency radio telescopes, such as the Low Frequency Array (LOFAR)\\footnote{http://www.lofar.org} and the Long Wavelength Array (LWA, \\eg Ellingson et al. 2009), might find in the upcoming years.  \\noindent LOFAR promises an impressive gain of two orders of magnitude in sensitivity and angular resolution over present instruments in the frequency range 15--240 MHz, and as such will open up a new observational window to the Universe. In particular, LOFAR is expected to contribute significantly to the understanding of the origin and evolution of the relativistic matter and magnetic fields in galaxy clusters.  The main focus of the present paper is to provide a theoretical framework for the interpretation of future LOFAR data by quantifying expectations for the properties and occurrence of giant radio halos in the context of the {\\it turbulent re-acceleration} scenario. In particular, in Sect. 2 we summarize the main ingredients used in the model calculations and provide an extension of the results of previous papers on the occurrence of radio halos in clusters (Sect.~2.1) and on the expected radio halo luminosity functions (Sect. 2.2). In Sect.~3 we derive the expected number counts of radio halos at 120 MHz and explore the potential of LOFAR surveys. Our conclusions are given in Sect.4.  \\noindent  A $\\Lambda$CDM ($H_{o}=70$ $\\mathrm{ Km\\, s^{-1} Mpc^{-1}}$, $\\Omega_{m}=0.3$, $\\Omega_{\\Lambda}=0.7$) cosmology is adopted throughout the paper. ",
        "conclusions": "In the present paper we perform Monte Carlo simulations to model the formation and evolution of giant radio halos in the framework of the merger-induced particle acceleration scenario (see Sec. 2). Following Cassano et al. (2006a) we use homogeneous models that assume {\\it a)} an average value of the magnetic field strength in the radio halo volume that scales with cluster mass as $B = B_{<M>} M_v^b$, and {\\it b)} that a fraction, $\\eta_t$, of the $P dV$ work, done by subclusters crossing the main clusters during mergers goes into {\\it magneto-acoustic} turbulence. Although simple, these models reproduce the presently observed fraction of galaxy clusters with radio halos and the scalings between the monochromatic radio power of halos at 1.4 GHz and the mass and X-ray luminosity of the host clusters (\\eg Cassano et al.~2006, 2008; Venturi et al. 2008), provided that the model parameters $(B_{<M>}, b , \\eta_t)$ lie within a fairly constrained range of values (Fig.~7 in Cassano et al. 2006a); in the present paper we adopt a reference set of parameters: $<B>=1.9\\, \\mu$G, $b=1.5$, $\\eta_t=0.2$, that falls in that range.   In Fig.~\\ref{fig:NC_NVSS_GMRT} we show that the expected number counts of giant radio halos at $\\nu_o=1.4$ GHz obtained with this set of parameters are nicely in agreement with both the data at low redshift (NVSS-XBACS selected radio halos, Giovannini et al.~1999) and at intermediate redshift (clusters in the ``GMRT radio halo survey'', Venturi et al.~2007, 2008).   The most important expectation of the {\\it turbulent re-acceleration} scenario is that the synchrotron spectrum of radio halos should become gradually steeper above a frequency, $\\nu_s$ that is determined by the energetics of the merger events that generate the halos and by the electron radiative losses (\\eg Fujita et al.~2003; Cassano \\& Brunetti 2005). Consequently, the population of radio halos is expected to be constituted by a mixture of halos with different spectra, with steep spectrum halos being more common in the Universe than those with flatter spectra (\\eg Cassano et al.~2006). The discovery of these very steep spectrum halos will allow to test the above theoretical conjectures.  In Sect.~2 we have derived the expected radio halo luminosity functions (RHLF) at frequency $\\nu_o$, that account for the contribution from the different populations of radio halos with $\\nu_s \\geq \\nu_o$. The RHLF are obtained combining the theoretical mass function of radio halos (with different $\\nu_s \\geq \\nu_o$) with the radio power--cluster mass correlation (Eq.\\ref{RHLF}). The expected monochromatic radio power at $\\nu_o$ of halos hosted by clusters with mass $M_v$ is extrapolated from the observed $P(1.4)$--$M_v$ correlation by assuming simple scaling relations, appropriate for homogenous models, that account for the dependence of the emitted synchrotron power on $\\nu_s$ (Eqs.\\ref{scale2},\\ref{scale3}).   \\noindent As a relevant case we calculate the expected RHLF at $\\nu_o$= 120 MHz (Fig.~\\ref{Fig.RHLF}). The shape of the RHLF can be approximated by a power law over more than two orders of magnitude in radio power. Homogeneous models imply the following scalings between $\\nu_s$, cluster mass and the radio luminosity at $\\nu_o$, $P_{\\nu_s} (\\nu_o)$ :  \\begin{equation} \\nu_s \\propto M^{4/3 +b} {{(1 + \\Delta M / M)^3}\\over {( <B>^2 + B_{cmb}^2 )^2}} \\label{nus_concl} \\end{equation}  \\noindent and from Eq.~\\ref{scale3} and the $P_{1.4}-M_v$ correlation:  \\begin{equation} P_{\\nu_s} (\\nu_o) \\propto M_v^3 \\nu_s^{\\alpha} \\label{p_concl} \\end{equation}  \\noindent i.e., radio halos with larger $\\nu_s$ are typically generated in massive clusters that undergo major mergers and contribute to the RHLF at larger powers. On the other hand halos with smaller $\\nu_s$ are typically generated in less massive systems and contribute to the RHLF at fainter powers. Radio halos with $\\nu_s \\geq$ 120 MHz however become increasingly rare in clusters with mass $\\leq 5 \\times 10^{14}$M$_{\\odot}$ and this explains the drop of the RHLF at lower radio powers in Fig.~\\ref{Fig.RHLF}. At the same time, halos with monochromatic radio emission at 120 MHz $> 10^{26}$W Hz$^{-1}$ would be generated in connection with very energetics merging events in very massive clusters, that are extremely rare, and this explains the RHLF cut-off at higher synchrotron powers in Fig.~\\ref{Fig.RHLF}.  In Sect.~3 we discuss the expected number counts of radio halos at 120 MHz, that best allow us to explore the potential of incoming LOFAR surveys in constraining present models.  \\noindent The crucial step in this analysis is the estimate of the minimum diffuse flux from giant radio halos that is detectable by these surveys. Because the LOFAR capabilities will become more clear during the incoming commissioning phase we exploit two complementary approaches: {\\it i)} we required that at least half of the radio halo emission is above a fixed brightness--threshold, $\\xi\\,F$ ($F$ being the $rms$ of LOFAR surveys; {\\it ii)} we required that the signal from the radio halo is $ \\geq 3 \\times \\xi\\,F$ in at least 5 beam area of LOFAR observations. In both cases we assume a fixed shape of the radial profile of radio halos, calibrated through that of several well studied halos at 1.4 GHz, that introduces a potential source of uncertainty.  Although the uncertainties due to the unavoidable simplifications in our calculations, the expected number counts of radio halos highlights the potential of future LOFAR surveys. By assuming the expected sensitivity of the LOFAR all sky survey (\\eg R\\\"ottgering 2009; priv. comm.), rms =0.1 mJy/b, and $\\xi \\sim 2-3$, we predict that about 350 giant radio halos ($\\sim 200$ considering the case {\\it ii)}) can be detected at redshift $\\leq$0.6. This means that LOFAR will increase the statistics of these sources by a factor of $\\sim 20$ with respect to present day surveys. About 55\\% of these halos are predicted with a synchrotron spectral index $\\alpha > 1.9$ between 250-600 MHz, and would brighten only at lower frequencies, unaccessible to present observations. Most important, the spectral properties of the population of radio halos are expected to change with increasing the sensitivity of the surveys as steep spectrum radio halos are expected to populate the low-power end of the RHLF. A large fraction of radio halos with spectrum steeper than $\\alpha \\approx 1.5$ (\\eg Fig.~\\ref{Fig.alpha}) is expected to allow a prompt discrimination between different models for the origin of radio halos, for instance in this case simple energetic arguments would rule out a secondary origin of the emitting electrons (\\eg Brunetti 2004; Brunetti et al. 2008).  Due to the large number of expected radio halos, a potential problem with these surveys is the identification of halos and of their hosting clusters. As a matter of fact we expect that LOFAR surveys will open the possibility to unveil radio halos in galaxy clusters with masses $\\gtsim 6-7 \\times 10^{14}$M$_{\\odot}$ at intermediate redshift. On the other hand, statistical samples of X-ray selected clusters, that are unique tools for the identification of the hosting clusters, typically select more massive clusters at intermediate z. Consequently, we explored the potential of the first LOFAR surveys as deep follow ups of available X-ray selected samples of galaxy clusters.  \\noindent We calculate the radio halo number counts expected from the follow up of clusters in the eBCS and MACS samples that collect $\\sim 400$ galaxy clusters in the redshift range 0--0.6. We expect that the LOFAR all sky survey, with a planned sensitivity in line with the case $\\xi\\,F=0.25$ mJy/b, will discover about 130 radio halos in eBCS and MACS clusters and that about 40\\% of these radio halos should have very steep spectrum, $\\nu_s \\leq 600$ MHz. The majority of radio halos in eBCS and MACS clusters are expected at z = 0.2-0.4, while the small number of clusters at $z \\geq 0.5$ in the MACS catalog does not allow a statistically solid expectations, although we may expect a couple of radio halos hosted in MACS clusters at this redshift.  The MS$^3$ survey will be carried out in 2010, covering the northern hemisphere, and is expected to reach a noise level of about 0.5 mJy/b at 150 MHz, that implies a sensitivity to diffuse emission from galaxy clusters about one order of magnitude (assuming $\\alpha \\approx 1.3$) better than present surveys (\\eg NVSS, Condon et al.~1998; VLSS, Cohen et al.~2007; WENSS, Rengelink et al.~1997).  \\noindent We considered MS$^3$ pointings relative to the fields of the about 300 galaxy clusters at $z\\leq0.3$ in the eBCS catalogues. We find that about 60 radio halos are expected to be detected by MS$^3$ observations in these clusters, 25\\% of them (10-15 halos) are expected with $\\nu_s \\leq$ 600 MHz. Fairly sensitive GMRT observations of eBCS clusters at redshift 0.2--0.3 are already available (Venturi et al.~2007, 2008) and we expect that in a few cases radio halos would glow up in the MS$^3$ images where no diffuse radio emission is detected at 610 MHz. We also find that MS$^3$ observations of eBCS clusters at $z=0-0.2$ can be used to test the increase of the fraction of cluster with radio halos with the X-ray luminosity of the host clusters, which is a unique prediction of our model (Fig.\\ref{Fig.histo_Lx}.   The most important simplification in our calculations is the use of homogeneous models. Non-homogeneous approaches, that model the spatial dependence of the acceleration efficiency and magnetic field in the halo volume (\\eg Brunetti et al. 2004), and possibly their combination with future numerical simulations, will provide a further step to interpret LOFAR data. Also the use of the extended PS theory is expected to introduce some biases. For instance, it is well-known that the PS mass function underproduces the expected number of massive clusters ($M>10^{15}\\,M_{\\odot}$) at higher redshift, $z\\sim 0.4-0.5$, by a factor of $\\sim 2$ with respect to that found in N-body simulations (\\eg Governato et al. 1999; Bode et al. 2001; Jenkins et al. 2001). Since in our model the great majority of halos at these redshift is associated with massive clusters, the use of the PS mass function implies that the RHNC at $z> 0.4-0.5$ could be underestimated by a similar factor. A further refinement of the approach proposed in the present paper could be obtained with the use of galaxy cluster {\\it merger trees} extracted from N-body simulations. These would also allow for a more {\\it realistic} description of the merger events (spatially resolved, multiple mergers, etc...).  In the present paper we focus on a reference set of model parameters. Cassano et al. (2006a) discussed the dependence of model expectations at 1.4 GHz on these parameters. Based on their analysis we expect that all the general results given in the present paper are independent of the adopted values for parameters. The expected number counts of halos should change by a factor of $\\sim2-2.5$ considering sets of model parameters within the region ($<B>$, $b$, $\\eta_t$) that allows for reproducing the observed $P_{1.4}-M_v$ correlation. In this case the number of halos that we expect decreases from super-linear sets of parameters ($b>1$ and $<B>\\geq 1.5\\,\\mu$G) to sub-linear cases ($b<1$ and $<B>\\leq 1.5\\,\\mu$G) (see also Fig.4 in Cassano et al. 2006b); a more detailed study will be presented in a future paper."
    },
    "0910/0910.5928_arXiv.txt": {
        "abstract": "The radii of some transiting extrasolar giant planets are larger than would be expected by the standard theory. We address this puzzle with the model of coupled radius-orbit tidal evolution developed by \\citet{Ibgui_and_Burrows_2009}. The planetary radius is evolved self-consistently with orbital parameters, under the influence of tidal torques and tidal dissipation in the interior of the planet. A general feature of this model, which we have previously demonstrated in the generic case, is that a possible transient inflation of the planetary radius can temporarily interrupt its standard monotonic shrinking and can lead to the inflated radii that we observe. In particular, a bloated planet with even a circular orbit may still be inflated due to an earlier episode of tidal heating. We have modified our model to include an orbital period dependence of the tidal dissipation factor in the star, $Q'_{\\ast} \\propto P^{\\gamma}$, $-1 \\leqslant \\gamma \\leqslant 1$. With this model, we search, for a tidally heated planet, orbital and radius evolutionary tracks that fall within the observational limits of the radius, the semimajor axis, and the eccentricity of the planet in its current estimated age range. We find that, for some inflated planets (WASP-6b and WASP-15b), there are such tracks; for another (TrES-4), there are none; and for still others (WASP-4b and WASP-12b), there are such tracks, but our model might imply that we are observing the planets at a special time. Finally, we stress that there is a two to three order-of-magnitude timescale uncertainty of the inspiraling phase of the planet into its host star, arising from uncertainties in the tidal dissipation factor in the star $Q'_{\\ast}$. ",
        "introduction": "\\label{sec:intro}  The more than 60 transiting extrasolar giant planets (EGPs) discovered so far\\footnote{ See J. Schneider's Extrasolar Planet Encyclopaedia at http://exoplanet.eu, the Geneva Search Programme at http://exoplanets.eu, and the Carnegie/California compilation at http://exoplanets.org.} offer a unique opportunity to test and improve the models of the structure and evolution of these bodies. The mass and radius of such planets can be inferred from a combination of radial velocity and transit lightcurve measurements that break the planet mass-inclination angle degeneracy. A large theoretical effort has been undertaken for more than a decade now to model and understand the evolution and the radii of transiting planets \\citep{guillot_et_al1996, burrows_et_al2000, Bodenheimer_et_al_2001, burrows_et_al2003, Baraffe_et_al_2003, Bodenheimer_et_al_2003,Gu_et_al_2003, burrows_et_al2004, Fortney_and_Hubbard_2004, Baraffe_et_al_2004, Chabrier_et_al_2004, laughlin_et_al_2005_1, Baraffe_et_al_2005, Baraffe_et_al_2006, burrows_et_al2007, Fortney_et_al_2007, Marley_et_al_2007, chabrier+baraffe2007, Liu_et_al_2008, Baraffe_et_al_2008, Ibgui_and_Burrows_2009, Miller_et_al_2009, Leconte_et_al_2009, Ibgui_et_al_2009_1}.  The radius of a gas giant planet depends on many physical effects that are particular to a given planet-star system, including the mass and age of the planet; the stellar irradiation flux and spectrum; the composition -- in particular, the heavy-element content -- of the atmosphere, the envelope, and the core; the atmospheric circulation that couples the day and the night sides; and any processes that could generate an extra power source in the interior of the planet, such as tidal heating. Moreover, the transit radius effect \\citep{burrows_et_al2003,Baraffe_et_al_2003} has to be considered in order to infer the transit radius from the planet's physical radius. Therefore, a custom evolutionary calculation is the most appropriate way to determine a theoretical transit radius, in order to compare it with the observed transit radius.  The objective of the present paper is to test the coupled radius-orbit tidal evolution model developed by \\citet{Ibgui_and_Burrows_2009} on some recently discovered inflated planets. We present evolutionary tracks for WASP-4b \\citep{Wilson_et_al_2008, Southworth_et_al_2009_2, Gillon_et_al_2009_1, Winn_et_al_2009_1} and WASP-12b \\citep{Hebb_et_al_2009}. We have also tested the model on TrES-4 \\citep{Mandushev_et_al_2007, Sozzetti_et_al_2008}, WASP-6b \\citep{Gillon_et_al_2009_2}, and WASP-15b \\citep{West_et_al_2009}. Some other inflated planets have been discovered recently, such as HAT-P-13b \\citep{Bakos_et_al_2009_2}, WASP-17b \\citep{Anderson_et_al_2009}, and CoRoT-5b \\citep{Rauer_et_al_2009}. We might apply our formalism to them in the future.  The idea of exploring tidal heating as an explanation for the inflated radii was originally formulated by \\citet{Bodenheimer_et_al_2001}. They suggested an excitation mechanism to sustain a nonzero eccentricity, for example a planetary companion \\citep{Bodenheimer_et_al_2003,Mardling_2007}. To date, two transiting EGPs are known to be accompanied by a companion, HAT-P-13b \\citep{Bakos_et_al_2009_2} and HAT-P-7b \\citep{Pal_et_al_2008,Winn_et_al_2009_4}. This reinforces the plausibility of such a scenario.  \\citet{Batygin_et_al_2009_2} have coupled a three-body tidal orbital evolution model with a model of the interior structure of HAT-P-13b. They found a quasi-stationary solution and possibly consistent core masses, radii, and tidal heating rates. Assuming, as did \\citet{Liu_et_al_2008}, that the systems are in a quasi-stationary state, \\citet{Ibgui_et_al_2009_1} provide, for each of the systems that they have studied (WASP-6b, WASP-12b, WASP-15b, TrES-4, HAT-P-12b) and for each of the associated atmospheric opacities (solar, $10\\times$solar) a relation between the heavy-element core mass $M_{\\rm core}$ and the ratio $e^2/Q'_{p}$, where $e$ is the orbital eccentricity and $Q'_{p}$ the tidal dissipation factor in the planet. This constraint results from a degeneracy between the dissipation heating rate in the interior of the planet (which increases the radius) and the mass of a possible heavy-element central core (which shrinks the radius).  For close-in EGPs (orbital separation $\\sles~0.1-0.15$~AU), tidal torques are strong enough such that they can result in planetary orbital evolution, and they produce tidal heating (dissipation) inside the planet. Such tidal effects were first suggested for transiting EGPs by \\citet{Jackson_et_al_2008_1,Jackson_et_al_2008_2, Jackson_et_al_2008_3}. Jackson et al. included the tides raised on the star and the tides raised on the planet. They found tidal rates close to the levels that \\citet{burrows_et_al2007} proposed to maintain the observed radii of some transiting EGPs. \\citet{Ibgui_and_Burrows_2009} described a model that couples the two consequences of these tidal effects -- planetary radius evolution and orbit evolution. They tested their model on HD~209458b and found an explanation for the radius of this planet. Note that they also showed that a supersolar metallicity of the planetary atmosphere, without invoking tides, can explain the radius. \\citet{Miller_et_al_2009} applied a similar method to all the transiting EGPs, albeit with simplified models for the atmospheres of the planets and parent stars, and a restricted range of possibilities for the tidal dissipation factors $Q'$. We have chosen a more detailed approach by adopting customized atmospheric models and an extended range for $Q'$. Therefore, due to the complexity of the atmospheric calculations, we have selected a couple of planets. We believe that both approaches are complementary. The one followed by \\citet{Miller_et_al_2009} provides a global estimate of the possibilities for matching the observed radii of the transiting EGPs, while our approach, more precise and therefore applied to fewer planetary systems, focuses on more specific issues such as the influence of the atmospheric opacity of the planet, a more detailed model for the tidal dissipation in the star, and a phenomenological study of all the qualitatively possible behaviors of the evolutionary curves \\citep{Ibgui_and_Burrows_2009}. The application of our model to a subset of inflated transiting EGPs is the subject of this paper.  The paper outlines, in Section~\\ref{sec:model}, the main assumptions of our coupled radius-orbit tidal evolution model, with a summary of the basic phenomenological results obtained from the previous study by \\citet{Ibgui_and_Burrows_2009}. It also explains our upgraded modeling of the tidal dissipation factor in the star. Section~\\ref{sec:M_Q_a_effect} demonstrates some additional generic results, such as the effect of $M_{\\rm core}$, $Q'_{p}$, and the initial semimajor axis. For each of the following planets, TrES-4, WASP-4b, WASP-6b, WASP-12b, and WASP-15b, we search for evolutionary tracks that fall within the observational limits of the radius, the semimajor axis, and the eccentricity of the planet in its current estimated age range. Our results are described in Section~\\ref{sec:applications}. We have not found any such track for TrES-4. We have found solutions for WASP-6b and WASP-15b. The cases of the planets WASP-4b and WASP-12b, for which we have coupled models that fit, are more interesting. Therefore, we present evolutionary curves for these planets. In fact, the solutions that we obtain for these two planets are valid only for very short age ranges in comparison with the estimated ages of the planets. This would imply that we are observing both planets at a very special time in their evolution, which would be a priori unlikely. In Section~\\ref{sec:plunging_timescale}, we discuss the plunging timescale of a planet into its host star. Its uncertainty can span two to three orders of magnitude. We summarize our results in Section~\\ref{sec:conclusion}.\\\\ ",
        "conclusions": "\\label{sec:conclusion}  We have presented in this paper some new general results of the coupled radius-orbit evolutionary model described in \\citet{Ibgui_and_Burrows_2009}, and we have applied the model to the inflated planets WASP-4b and WASP-12b. We assumed a two-body gravitational and tidal interaction between the planet and its host star, coupling the planetary radius and the orbit evolution. We included the tides raised on the planet and the tides raised on the star. Stellar irradiation and a detailed planetary atmosphere are included. The fundamental result is the transient inflation of the planetry radius that temporarily interrupts its monotonic standard shrinking. An important point is that even though the current orbit of the planet has almost circularized, the radius of the planet can still be inflated due to an earlier episode of tidal heating. This is why we stress that an inflated planet with an observed circular orbit can still have tidal heating as an explanation of its radius. Fixing the planet and star properties, the model is controlled by four free parameters, ($Q'_{p},Q'_{\\ast},e_{i},a_{i}$), that are the tidal dissipation factors in the planet and in the star, and the initial eccentricity and semimajor axis at the beginning of this two-body evolution. We stress the sensitive and nonlinear dependence of the evolutionary curves on these parameters.  We have demonstrated that an increase of either the core mass $M_{\\rm core}$, or $Q'_{p}$, or $a_{i}$ results in a lower value of the radius inflation peak and in a delay of its appearance.  The final semimajor axis is the same, whatever $M_{\\rm core}$ or $Q'_{p}$, but is larger when $a_{i}$ is larger.  At an earlier age, the planet with the larger core has the smaller radius, but this is opposite at later ages.  We have enhanced our model by including an orbital period dependence of the tidal dissipation in the star, $Q'_{\\ast} \\propto P^{\\gamma}$, $-1 \\leqslant \\gamma \\leqslant 1$.  $Q'_{\\ast}$ drives the inspiral of the planet into its host star.  Applications of our model to recently detected transiting inflated planets show that: \\begin{itemize}\\itemsep-0.04in% \\vspace{-6pt} \\item WASP-6b and WASP-15b can be fit at solar opacity over Gyr age ranges. \\item We have not found an acceptable fit for TrES-4, at either solar, $3\\times$solar, or $10\\times$solar planet atmospheric opacity. \\item WASP-4b can be fit at solar opacity with, for example, the combination ($Q'_{p},e_{i},a_{i}$)=($10^{8.0},0.80,0.050$) and with $Q'_{\\ast}=10^{6.5}\\times(P/\\rm 10days)^{-1}$. \\item WASP-12b can be fit at solar opacity with, for example, the combination ($Q'_{p},e_{i},a_{i}$)=($10^{8.0},0.73,0.055$) and with $Q'_{\\ast}=10^{6.5}\\times(P/\\rm 10days)^{-1}$. \\end{itemize} \\vspace{-6pt} For WASP-4b and WASP-12b, the ranges of ages that allow simultaneous fits of radius, semimajor axis, and orbital eccentricity, are very narrow, seeming to suggest that, if the two-body coupled evolutionary model described herein is in fact responsible for these planets' inflated radii, then we are observing them at a special epoch in their evolution.  However, relaxing the fit-criterion from 1$\\sigma$ to 2$\\sigma$ or 3$\\sigma$ would alleviate this apparent problem.  Our results (in particular, for TrES-4) suggest that a coupled radius-orbit tidal evolution model might not on its own explain the radii of all the inflated transiting giant planets.  An alternative scenario with stationary heating has been proposed \\citep{Ibgui_et_al_2009_1} and applied to all the planets discussed in this paper. Though not providing direct solutions to the inflated radii issue, this scenario constrains the ratio $e^{2}/Q'_{p}$ for a given $M_{\\rm core}$. Finally, a combination of these two models could be imagined with a two-body interaction, followed by a quasi-steady low eccentricity phase due to perturbations by a second planet.  The last point we make in this paper is the uncertainty of the plunging timescale during the spiraling of the planet into its host star. This timescale is strongly dependent on the semimajor axis; specifically, it depends on $a$ to the 6.5 power. It also has a linear dependence on $Q'_{\\ast}$, which is a parameter that is uncertain by several orders of magnitude. We have shown that HD~209458b can plunge in between 0.5 and 60~Gyr from now, a 2-order-of-magnitude range, and that WASP-12b can plunge in between 0.1 and 100~Myr from now, a 3-order-of-magnitude range.  \\citet{Ibgui_and_Burrows_2009} have suggested caveats to, and ways to improve, the model employed here. We close by noting several additional points. The orbital evolution equations depend on the theory of tidal dissipation inside gaseous planets and stars. Improvements to this theory might result in different evolutionary tracks (of $R_p$, $e$, and $a$) from the ones presented in this paper. Furthermore, we have noted that it is not strictly appropriate to apply these equations to model scenarios with large values of orbital eccentricity.  However, both for comparison with previous work and because the proper tidal theory remains unknown, our present approach is a valuable step in exploring the extent to which tidal dissipation might explain the radii of the inflated EGPs. Further observations that might help to constrain this model and to discriminate between this and the stationary-state model of \\citet{Ibgui_et_al_2009_1} include both increasing the accuracy of orbital eccentricity measurements and searching for companions to the transiting EGPs. These and other advances will help us progress toward a better understanding of the puzzle of the inflated planets."
    },
    "0910/0910.0399_arXiv.txt": {
        "abstract": "Sagittarius\\,A$^\\star$ (\\sgra) is the supermassive black hole residing at the center of the Milky Way. It has been the main target of an extensive multiwavelength campaign we carried out in April 2007. Herein, we report the detection of a bright flare from the vicinity of the horizon, observed simultaneously in X-rays (\\xmm/EPIC) and near infrared (\\vlt/NACO) on April 4$^{\\rm th}$ for 1--2~h. For the first time, such an event also benefitted from a soft $\\gamma$-rays (\\integ/ISGRI) and mid infrared (\\vlt/VISIR) coverage, which enabled us to derive upper limits at both ends of the flare spectral energy distribution (SED). We discuss the physical implications of the contemporaneous light curves as well as the SED, in terms of synchrotron, synchrotron self-Compton and external Compton emission processes. ",
        "introduction": "From the discovery of a compact radio source, \\sgra, at the Galactic Center (GC) in 1974 \\citep{balick74} to the near infrared (NIR) tracking of stars in Keplerian motion around \\sgra~three decades later \\citep{schodel02,ghez03}, the evidence for a $\\sim$$4\\times10^{6}$~$M_\\odot$ black hole with very slow proper motion at the dynamical center of our galaxy \\citep{reid08} gradually piled up \\citep[see][for a general review and references therein]{melia07}.  Yet, the long quest for the high energy emission pertaining to the black hole has only been achieved recently. \\sgra~was resolved as a notably dim ($2.4\\times10^{33}$ \\ergpersec, 2--10~keV) and slightly extended (1.4$''$) point source with the \\chandra~satellite in 1999 \\citep{baganoff03a}. One year later, the same instrument witnessed the source exhibiting an X-ray flare for $\\sim$3 h \\citep{baganoff01}. A $\\sim$10~min long substructure within the light curve of the eruption and light time travel arguments imply that this event took place close to the  event horizon ($<15~R_{\\rm S}$). Many other detections of X-ray flares followed, either with \\xmm~or \\chandra~\\citep[see e.g.][]{goldwurm03a,baganoff03b,porquet03,belanger05}, and established that the duty cycle of the black hole is nearly one X-ray flare per day. The origin of these events is still unclear, in spite of all the efforts aimed at their monitoring in different energy ranges. In 2003, NIR flares from \\sgra~were indeed discovered with the \\vlt~\\citep{genzel03}, and later confirmed by the \\keck~\\citep{ghez04} and the \\hst~\\citep{yusef-zadeh06a}. They occur more frequently than the X-ray ones (around four per day) and have been observed in many NIR atmospheric pass bands (H, K, L, M). Each new infrared flare has generally induced either spatial \\citep{clenet05}, spectral \\citep{eisenhauer05,ghez05,gillessen06,krabbe06,hornstein07}, polarimetric \\citep{eckart06b,meyer06,meyer07,trippe07}, or timing studies \\citep{meyer08,do09}. Numerous multiwavelength campaigns showed that an X-ray flare always comes along with a  simultaneous NIR one\\footnote{The converse is not true, some NIR flares have no X-ray counterpart \\citep{hornstein07}.} \\citep{eckart04,eckart06a,eckart08a,yusef-zadeh06a,hornstein07}, and maybe a delayed submm one \\citep{marrone08,eckart08b,yusef-zadeh08} caused by plasmon expansion \\citep{liu04,yusef-zadeh06b}.  Above 6~$\\mu$m, in the mid infrared (MIR), no detection of \\sgra~has been reported so far. Recent upper limits on the black hole flux at 8.6~$\\mu$m were set by the \\vlt/VISIR instrument during low level NIR variability by \\citet{schodel07}, who argued that a detection would be reachable in case of a strong NIR flare.  Above 20 keV, repeated surveys of the heart of the Milky Way in soft $\\gamma$-rays with the \\integ~satellite unveiled a persistent pointlike source compatible with \\sgra~location (within the 1$'$ error radius), \\igr~\\citep{belanger04,belanger06}. The nature of the source is still uncertain, and a possible association with the supermassive black hole remains conceivable. Given the limited angular resolution of the soft $\\gamma$-ray telescope \\integ/IBIS/ISGRI ($\\sim$12$'$ FWHM), the best way to unequivocally identify the mysterious \\igr~with \\sgra~is the detection of correlated variability between soft $\\gamma$-rays and other wavelengths.  To tackle the above puzzles and investigate the correlated X-ray/NIR variability of \\sgra~in more details, a coordinated multiwavelength campaign on the GC was conducted in spring 2007. It involved in particular the \\xmm~and \\integ~satellites for the high energies, as well as the \\vlt/ NACO and \\vlt/VISIR ground instruments to cover the NIR and MIR part of the spectrum, respectively. Their results are presented in Sect.~\\ref{obs} and interpreted in Sect.~\\ref{SED}. Note that the X-ray and infrared findings have already been published by \\citet{porquet08} and \\citet{dodds09}, respectively.  We will not discuss here the short term variability of \\sgra~in April 2007 at cm, mm, and submm wavelengths, which will be reported in another article, along with NIR results obtained by the Hubble Space Telescope \\citep{yusef-zadeh09}\\footnote{For further discussion of the past variability of \\sgra~in cm and mm bands, see for example \\citet{zhao01,herrnstein04} and \\citet{tsuboi99,zhao03}, respectively.}.  Throughout this paper we adopt a GC distance of 8 kpc \\citep{reid93} and a black hole mass $M_{\\bullet}=4\\times10^{6}$~$M_\\odot$ \\citep{ghez08}, for which the Schwarzschild radius is $R_{\\rm S}=1.2\\times10^{12}$~cm.   \\begin{figure*}  \\centering \\includegraphics[trim=1cm 2cm 1cm 1.5cm, clip=true, width=14cm]{Fig1.pdf} \\caption{Multiwavelength and multiscale views of the Galactic Center in April 2007 in R.A. ($^\\circ$) horizontally and Dec. ($^\\circ$) vertically (North and East point towards the top and the left, respectively). \\emph{Top, left}: INTEGRAL/ISGRI mosaic  in the 20--40~keV band. The Galactic plane runs from upper left to bottom right. \\emph{Top, right}: INTEGRAL/JEM-X 1 mosaic  in the 3--20~keV band. \\emph{Middle}: XMM/PN image in the 2--10~keV band. \\emph{Bottom, left}: VLT/NACO image at the peak of the flaring period at 3.8~$\\mu$m. \\emph{Bottom, right}: VLT/VISIR average image of the 3$^{\\rm rd}$/4$^{\\rm th}$ April night at 11.88~$\\mu$m. The three dusty arms swirling around \\sgra~compose the so-called ``mini-spiral''.} \\label{ima} \\end{figure*} ",
        "conclusions": "This paper complements a series of articles about the April 2007 synchronous observations of the Galactic Center from radio to $\\gamma$-rays \\citep{porquet08,dodds09,yusef-zadeh09}. Here, we have recapped the results on the brightest flare ever detected simultaneously at NIR and X-ray frequencies. We have also reported for the first time $\\gamma$-ray constraints on such an event, which, added to our MIR/NIR/X-ray spectral measurements, constitute the broadest simulaneous spectrum of a flare ever achieved. The essential observational conclusions may be summarized as follows: \\begin{itemize} \\item the peaks of the X-ray and NIR emissions are coincident within 3~min; \\item the width of the NIR flare light curve is broader than the X-ray one by a factor $\\sim$2; \\item the NIR light curve is substructured on a timescale of $\\sim$20~min while the X-ray light curve is rather smooth; \\item there is no detectable MIR counterpart; \\item the soft $\\gamma$-ray source \\igr~is non variable. \\end{itemize} The high quality of the spectral information we gathered allowed for a discussion of the several classical radiative processes models employed to explain the flares: SSC, EC, and SB. Yet, none of these mechanisms is entirely satisfactory to meet our observations. The theoretical inquiries to come will have to take into account the time evolution of the phenomenon and the aging of the radiating particles to better connect the light curves and spectra. From an observational stand point, it will be useful to repeat such NIR/X-ray measurements in a near future to get two respective individual and fully contemporaneous spectra, which has never been accomplished thus far. As we have seen, one key probe of what powers the flares, is a better determination of the X-ray spectral slope. In a more distant future, thanks to a broad X-ray sensitivity over the 1--80~keV band and a high angular resolution above 10~keV, \\simx~should address this issue and resolve the GC region in soft $\\gamma$-rays."
    },
    "0910/0910.5082_arXiv.txt": {
        "abstract": "{ The recently discovered planetary system HD45364 which consists of a Jupiter and Saturn mass planet is very likely in a 3:2 mean motion resonance. The standard scenario to form planetary commensurabilities is convergent migration of two planets embedded in a protoplanetary disc. When the planets are initially separated by a period ratio larger than two, convergent migration will most likely lead to a very stable 2:1 resonance for moderate migration rates. To avoid this fate, formation of the planets close enough to prevent this resonance may be proposed. However, such a simultaneous formation of the planets within a small annulus, seems to be very unlikely.  Rapid type III migration of the outer planet crossing the 2:1 resonance is one possible way around this problem. In this paper, we investigate this idea in detail. We present an estimate for the required convergent migration rate and confirm this with N-body and hydrodynamical simulations. If the dynamical history of the planetary system had a phase of rapid inward migration that forms a resonant configuration, we predict that the orbital parameters of the two planets are always very similar and hence should show evidence of that.  We use the orbital parameters from our simulation to calculate a radial velocity curve and compare it to observations. Our model can explain the observational data as good as the previously reported fit. The eccentricities of both planets are considerably smaller and the libration pattern is different. Within a few years, it will be possible to observe the planet-planet interaction directly and thus distinguish between these different dynamical states. }  \\date{Submitted: 30 August 2009 - Revised: 26 October 2009 - Accepted: 26 October 2009 } ",
        "introduction": "Over 400 extrasolar planets have already been discovered \\citep{exoplanet} and their diversity keeps challenging planet formation theory. For example, the recently discovered multi-planetary system HD45364 raises interesting questions about its formation history.  The planets have masses of $m_1=0.1906M_{\\text{Jup}}$ and $m_2=0.6891M_{\\text{Jup}}$ and are orbiting the star at a distance of $a_1=0.6813~\\text{AU}$ and $a_2=0.8972~\\text{AU}$, respectively \\citep{CorreiaUdry2008}. The period ratio is close to $1.5$ and a stability analysis implies that the planets are deep inside a 3:2 mean motion resonance. The planets have most likely formed further out in cooler regions of the proto-stellar disc, as water ice which is an important ingredient for dust aggregation can only exist beyond the snow line which is generally assumed to be at radii larger than $2~\\text{AU}$ \\citep{SasselovLecar2000}.  It is then usually assumed that migration due to planet disc interactions has moved the planets closer to the star. Although the details of this process are still hotly debated, the existence of many resonant multiplanetary systems supports this idea. During migration the planets can get locked into a commensurabilities after which the planets migrate together with a constant period ratio. In such a resonance, one or more resonant angles are librating \\citep[see e.g.][]{LeePeale01}.  For the planetary system HD45364 this standard picture poses a new problem. Assuming that the planets have formed far apart from each other, the outcome for the observed masses after migration is almost always a 2:1 mean motion resonance, not 3:2 as observed. The 2:1 resonance that forms is found to be extremely stable. One possible way around this is a very rapid convergent migration phase that passes quickly through the 2:1 resonance.  In this paper we explore this idea quantitatively. The plan of the paper is as follows. In section \\ref{sec:formationnbody} we show, using N-body simulations, that the system always ends up in the 2:1 resonance assuming moderate migration rates. We present scenarios which result in formation of a 3:2 resonance after a rapid migration phase. In section \\ref{sec:formationhydro} we perform hydrodynamic simulations with a variety of disk models in order to explore the dependence on the physical setup. In section \\ref{sec:otherformationscenarios} we briefly discuss other formation scenarios. We go on to compare the orbital parameters of our simulations with the observed radial velocity data in section \\ref{sec:observation}.  We find that the orbital parameters observed in our simulations differ from those estimated from a statistical fit by \\cite{CorreiaUdry2008}. However, our models reproduce the observational data very well and we argue that our fit has the same level of significance as that obtained by \\cite{CorreiaUdry2008}.   Future observations will be able to resolve this issue. This is the first prediction of orbital parameters for a specific extrasolar planetary system derived from planet migration theory only. The parameter space of orbital configurations produced by planet disc interactions (low eccentricities, relatively small libration amplitudes) is very small. As in case of the GJ876 system, this can provide strong evidence on how the system formed. Finally, we summarise our results in section \\ref{sec:conclusion}. ",
        "conclusions": "The planets in the multi-planetary system HD45364 are most likely in a 3:2 mean motion resonance. This poses interesting questions on it's formation history. Assuming that the planets form far apart from each other and migrate with moderate migration rate, as predicted by standard planet formation and migration theories, the most likely outcome is a 2:1 mean motion resonance, contrary to the observation of a 3:2 MMR.  In this work, we investigated a possible way around this problem by letting the outer planet undergo a rapid inward type~III migration. We presented an analytical estimate and performed N-body as well as hydrodynamical simulations.  We found that it is indeed possible to form a 3:2 MMR and avoid the 2:1 resonance, thus resembling the observed planetary system using reasonable disc parameters. Hydrodynamical simulations suggest that the system is more likely to sustain the resonance for large aspect ratios, as the migration of the inner planet is slowed down, thus avoiding divergent migration.  Finally, we used the orbital configuration found in the hydrodynamical formation scenario to calculate a radial velocity curve. This curve has then been compared to observations and the resulting fit has an identical $\\chi^2$ value than the previously reported \\textit{best fit}.  Our solution is stable for at least a million years. It is in a dynamically different state, both planets having lower eccentricities and a different libration pattern. This is the first time that planet migration theory can predict a precise orbital configuration of a multiplanetary system.  This might also be the first direct evidence for type III migration if this scenario turns out to be true. The system HD45364 remains an interesting object for observers, as the differences between the two solutions can be measured in radial velocity within a couple of years.     \\label{sec:conclusion}"
    },
    "0910/0910.0458_arXiv.txt": {
        "abstract": "The timescale for energy release is an important parameter for constraining the coronal heating mechanism. Observations of ``warm'' coronal loops ($\\sim 1$\\,MK) have indicated that the heating is impulsive and that coronal plasma is far from equilibrium.  In contrast, observations at higher temperatures ($\\sim 3$\\,MK) have generally been consistent with steady heating models.  Previous observations, however, have not been able to exclude the possibility that the high temperature loops are actually composed of many small scale threads that are in various stages of heating and cooling and only appear to be in equilibrium.  With new observations from the EUV Imaging Spectrometer (EIS) and X-ray Telescope (XRT) on \\textit{Hinode} we have the ability to investigate the properties of high temperature coronal plasma in extraordinary detail. We examine the emission in the core of an active region and find three independent lines of evidence for steady heating. We find that the emission observed in XRT is generally steady for hours, with a fluctuation level of approximately 15\\% in an individual pixel. Short-lived impulsive heating events are observed, but they appear to be unrelated to the steady emission that dominates the active region.  Furthermore, we find no evidence for warm emission that is spatially correlated with the hot emission, as would be expected if the high temperature loops are the result of impulsive heating. Finally, we also find that intensities in the ``moss,'' the footpoints of high temperature loops, are consistent with steady heating models provided that we account for the local expansion of the loop from the base of the transition region to the corona. In combination, these results provide strong evidence that the heating in the core of an active region is effectively steady, that is, the time between heating events is short relative to the relevant radiative and conductive cooling times. ",
        "introduction": "The coronal heating problem is one of the most fundamental open questions in solar physics. The formation of the high temperature solar atmosphere is clearly related to the magnetic fields generated in the solar interior, but how this magnetic energy is converted to the thermal energy of the corona remains unknown. One important constraint on the coronal heating mechanism is the time scale for energy release. If energy is deposited into magnetic flux tubes on timescales that are very short compared to a characteristic cooling time, then the corona will be filled with loops that are close to equilibrium and appear to be steady. Active region observations have generally suggested that high temperature loops are consistent with steady heating \\citep[e.g.,][]{porter1995,kano1995}. These studies have found that the observed evolution of the emission is much slower than the radiative and conductive cooling timescales. There has also been some success in modeling entire active regions with steady heating models \\citep{schrijver2004,warren2006b,winebarger2008,lundquist2008}. Finally, \\cite{antiochos2003} have argued that observations of the ``moss,'' the footpoints of high temperature loops, are also consistent with steady heating. They found that the average moss intensities are typically constant over many hours and loops cooling through 1\\,MK were not observed in the moss region they observed.  There is considerable evidence that coronal loops observed at lower temperatures ($\\sim 1$\\,MK) are evolving and not in equilibrium \\citep[e.g.,][]{aschwanden2001b,winebarger2003b,cirtain2007,urra2009}. These ``warm'' loops appear to have apex densities that are much higher than can be accounted for by steady heating models \\citep[e.g.,][]{winebarger2003,aschwanden2008b}. The properties of the warm loops are more consistent with impulsive heating models \\citep[e.g.,][]{warren2002b,spadaro2003,warren2003}.  If it is true that hot loops are close to equilibrium while the warm loops are generally heated impulsively then the coronal heating mechanism becomes even more difficult to understand. The cooling time for short, hot loops is relatively rapid (typically a few hundred seconds), so heating events on these loops would need to occur very frequently. The cooling time for the warm loops is much longer (typically a few thousand seconds), so impulsive heating events on these loops would be infrequent.  It is tempting to conjecture that the emission at high temperatures is also consistent with impulsive heating models and that the apparent steadiness of this emission is the result of the superposition of many evolving stands along the line of sight \\citep[e.g.][]{cargill1997,cargill2004,patsourakos2006}. Previous observations have not been able to exclude this possibility. For example, it could be that the cooling loops are not easy to detect in an active region core because they are faint relative to both the bright moss and the extended corona in which the active region is embedded.  The EUV Imaging Spectrometer (EIS) and the X-ray Telescope (XRT) on the \\textit{Hinode} mission provide a new opportunity to observe active region emission in unprecedented detail. EIS observes emission covering a very broad range of coronal temperatures: between \\ion{Fe}{7} and \\ion{Fe}{24}, only 4 ionization stages of Fe are not present in the EIS data (\\ion{Fe}{18}--\\textsc{xxi}). EIS also observes \\ion{Ca}{14}, \\textsc{xv}, \\textsc{xvi}, and \\textsc{xvii}, providing excellent coverage of the critical temperature range around 3\\,MK \\citep{warren2008c}. XRT is a broadband imaging telescope that observes high temperature plasma very efficiently and at high spatial resolution. The high cadence XRT data complement EIS, which often observes at a much lower cadence, and allows us to track the evolution of coronal plasma over a large field of view.  In this paper we use EIS and XRT observations to examine the properties of coronal plasma in the core of an active region. The active region that we have selected (NOAA active region 10960) is unusual in that most of the overlying warm loops are located to the north and south of the region, providing a largely unobstructed view of the active region core over a wide range of temperatures (see Figure~\\ref{fig:context}). High cadence EIS observations of this region, which was observed by \\textit{Hinode} during the period 4--13 June 2007, have shown that the \\ion{Fe}{12} 195.119\\,\\AA\\ intensities, Doppler shifts, and non-thermal widths in the moss are constant over long periods of time, suggesting steady heating \\citep{brooks2009b}. In this study we find three additional lines of evidence that also indicate that the heating of high temperature loops is steady. We examine the evolution of the emission in individual pixels in XRT and find that the vast majority of the emission is constant to within approximately 15\\% during this period. The inspection of emission at different temperatures in the core of the active region shows that there is no evidence for loops cooling from high temperatures. We find no relationship between the warm emission (\\ion{Fe}{10}--\\ion{Fe}{14}) and the steady hot emission (\\ion{Ca}{14}--\\ion{Ca}{17}). Finally, we find that we can bring all of the observed moss intensities into agreement with steady heating models, if we allow for loop constriction at the base of the loop. Previous modeling work had found significant discrepancies at the lowest temperatures observed with EIS \\citep{warren2008}.  None of these observations is conclusive; they do not necessarily exclude alternative models that involve non-equilibrium processes. However, these observations do provide very strong constraints on the mechanism responsible for producing the high temperature emission observed in solar active regions. ",
        "conclusions": "We have presented a comprehensive analysis of observations in the core of an active region using data from the EIS and XRT instruments on \\textit{Hinode} and \\textit{TRACE}. The apparent steadiness of the XRT emission, the lack of spatial correlation between the hot and warm emission, and the consistency of the funnel models with the observed emission all point to frequent heating events that keep the hot loops close to equilibrium. Furthermore, these results are consistent with high cadence EIS measurements of moss intensities, Doppler shifts, and nonthermal widths that show little evidence of dynamical events over many hours \\citep{brooks2009b}. In combination, these results provide strong evidence that the heating in the core of an active region is effectively steady, that is, the time between heating events is short relative to the relevant radiative and conductive cooling times."
    },
    "0910/0910.2946_arXiv.txt": {
        "abstract": "A modern Q-band low noise amplifier (LNA) front-end is being fitted to the 10.4~m millimeter-wave telescope at the Raman Research Institute (RRI) to support observations in the 40-50~GHz frequency range. To assess the suitability of the surface for this purpose, we measured the deviations of the primary surface from an ideal paraboloid using radio holography. We used the 11.6996 GHz beacon signal from the GSAT3 satellite, a 1.2~m reference antenna, commercial K$_{u}$-band Low Noise Block Convereters (LNBC) as the receiver front-ends and a Stanford Research Systems (SRS) lock-in amplifier as the backend. The LNBCs had independent free-running first local oscillators (LO). Yet, we recovered the correlation by using a radiatively injected common tone that served as the second local oscillator. With this setup, we mapped the surface deviations on a $64 \\times 64$ grid and measured an rms surface deviation of $\\sim 350~\\mu$m with a measurement accuracy of $\\sim 50~\\mu$m. ",
        "introduction": "RRI has a mm-wave Leighton telescope, of 10.4~m diameter with 81 hexagonal panels~\\citep{ref:tks}. It is being rejuvenated to undertake Q-band observations at 43~GHz. This requires the surface rms error to be below $\\sim 370~\\mu$m in order to have an aperture efficieny better than 50\\%, as given by Ruze's relation:~$\\eta_{ap}=\\eta_0 ~exp{(\\frac{4\\pi \\delta}{\\lambda})^2}$~\\citep{ref:ruze66}. Radio holographic surface measurement was carried out in Aug-Sep 2007 to measure the surface rms error, and to identify panels that may need correction. In this poster paper, we report the details of this experiment, analyse the data, discuss the results and present our conclusions.  \\begin{figure}[!t] \\centering \\epsfig{figure=holoRx4.eps, width=0.93\\linewidth} \\caption{\\small Dual channel holography receiver layout\\normalsize } \\label{f:receiver} \\end{figure} ",
        "conclusions": "The effect of free runnning local oscillators in LNBC was solved by radiatively injecting a tone. The effect of satellite drifting was removed by pointing to the satellite every 30~minutes. An SNR $>$ 5000 was achieved by making the receiver chain robust, repeating and co-adding the central block scans \\& using satellite pointing data for amplitude and phase calibration.  From the measured surface deviations (see Fig.~\\ref{f:map}(a)), the surface rms error is calculated to be $\\sim350~\\mu$m. It is within $\\lambda$/16, implying that Q Band observations are possible. It can also be seen that some panels require correction. Fig.~\\ref{f:map}(b) shows the residual obtained by subtracting two independent surface deviation measurements. The rms of this residual is $\\sim70~\\mu$m. Therefore, the measurement accuracy is estimated to be $\\sim50~\\mu$m."
    },
    "0910/0910.3037_arXiv.txt": {
        "abstract": "Based on the data obtained from the \\textit{Spitzer}/GLIPMSE Legacy Program and the 2MASS project, we derive the extinction in the four IRAC bands, [3.6], [4.5], [5.8] and [8.0]$\\mum$, relative to the 2MASS $\\Ks$ band (at 2.16$\\mum$) for 131 GLIPMSE fields along the Galactic plane within $|l|\\leq65^{\\rm o}$, using red giants and red clump giants as tracers. As a whole, the mean extinction in the IRAC bands (normalized to the 2MASS $\\Ks$ band), $A_{[3.6]}/A_\\Ks\\approx0.63\\pm0.01$, $A_{[4.5]}/A_\\Ks\\approx0.57\\pm0.03$, $A_{[5.8]}/A_\\Ks\\approx0.49\\pm0.03$, $A_{[8.0]}/A_\\Ks\\approx0.55\\pm0.03$, exhibits little variation with wavelength (i.e. the extinction is somewhat flat or gray). This is consistent with previous studies and agrees with that predicted from the standard interstellar grain model for $R_V=5.5$ by \\citet{Weingartner01}. As far as individual sightline is concerned, however, the wavelength dependence of the mid-infrared interstellar extinction $A_{\\lambda}/A_\\Ks$ varies from one sightline to another, suggesting that there may not exist a ``universal'' IR extinction law. We, for the first time, demonstrate the existence of systematic variations of extinction with Galactic longitude which appears to correlate with the locations of spiral arms as well as with the variation of the far infrared luminosity of interstellar dust. ",
        "introduction": "With the development of space infrared (IR) astronomy, the precise determination of IR extinction becomes urgent in order to recover the intrinsic colors and spectral energy distributions (SEDs) of heavily obscured sources. There have been various attempts to measure the IR extinction based on the \\emph{Infrared Space Observatory}(\\emph{ISO}) and \\emph{Spitzer Space Telescope} since \\citet{Lutz96} obtained the mid-IR extinction from several hydrogen recombination lines and demonstrated the absence of the model-predicted pronounced minimum around 7$\\mum$. This was supported by \\citet{Jiang03, Jiang06} based on the ISOGAL database \\citep{Omont99}, and by \\citet{Indebetouw05} based on the data from the \\emph{Spitzer} Galactic Legacy Infrared Midplane Survey Extraordinaire (GLIMPSE) Legacy Program \\citep{Benjamin03}. All these results roughly agree with the extinction predicted by the standard interstellar grain model for $R_V=5.5$ of \\citet{Weingartner01} and \\citet{Draine03}.\\footnote{% $R_V\\equiv A_V/E(B-V)$ is the total-to-selective extinction ratio, where $E(B-V)\\equiv A_B-A_V$, the color excess, is the difference between the extinction in $B$ and $V$ bands. } However, so far only a few wave bands have been investigated and the sky coverage is also limited. No consensus has been reached yet regarding the interstellar extinction at $\\simali$5--8$\\mum$.  Recent progress in the IR extinction measurements is made toward star-forming regions mainly based on the \\emph{Spitzer} observations. Thanks to the high sensitivity of \\emph{Spitzer}, deep photometry is now possible and objects that suffer severe extinction are now reachable. \\citet{Flaherty07} studied five nearby star-forming regions at mid-IR wavelengths (3.6$\\mum$, 4.5$\\mum$, 5.8$\\mum$ and 8.0$\\mum$ from the InfraRed Array Camera [IRAC],\\footnote{% The effective wavelengths of the four IRAC bands are actually 3.545$\\mum$, 4.442$\\mum$, 5.675$\\mum$ and 7.760$\\mum$, respectively. } and 24$\\mum$ from the Multiband Imaging Photometer [MIPS]). They confirmed a relatively flat extinction curve at $\\simali$4--8$\\mum$. \\citet{Roman07} studied a star-forming dense cloud core located in the Pipe Nebula, and found that the IR extinction in the IRAC bands of that region also agrees with the $R_{V}=5.5$ model curve and indicates a dust size distribution favoring larger sizes.  Although both the \\emph{ISO} and \\emph{Spitzer} measurements agree with each other in that the $\\simali$5--8$\\mum$ extinction is relatively flat and lacks the model-predicted minimum around 7$\\mum$ of \\citet{Draine89}, there do exist differences among various measurements made for different sightlines. In their sample of five sightlines toward star-forming regions, \\citet{Flaherty07} found a clear difference between one sightline and the other four sightlines (see their Table 3). They derived higher $A_{\\lambda}/A_\\Ks$ ratios and a flatter wavelength dependence than that of \\citet{Indebetouw05} for the same sightline toward $l=284^{\\circ}$ in the Galactic Plane.  From $>$\\,200 fields observed in the ISOGAL survey, \\citet{Jiang06} analyzed the extinction at 7$\\mum$ and 15$\\mum$ along $\\simali$120 directions. They found marginal variation of the extinction at 7$\\mum$. It is commonly believed that, with the parameter $R_V$ increasing in denser regions, the variation of the ultraviolet (UV) and visual extinction with wavelength becomes flatter than that of the diffuse interstellar medium (ISM) which is characterized with a lower $R_V$ \\citep{Cardelli89}, while the near-IR extinction seems to be ``universal\", with little variation among different sightlines \\citep{Draine89}. However, \\citet{Nishiyama06a, Nishiyama09} recently argued against such a universal near-IR extinction. In addition, \\citet{Fitzpatrick04} argued that the IR-through-UV Galactic extinction curves should not be considered as a simple one-parameter family, whether characterized by $R_V$ as suggested by \\citet{Cardelli89} or any other parameters.  \\citet{Whittet77} presented observational evidence for a small but appreciable variation in $R_V$ with Galactic longitude. He suggested that the most likely explanation for this is a variation in the mean size of the dust in the local spiral arm. However, unfortunately only a few data points were used in that work and therefore no systematic variation of the extinction with Galactic longitude was reported. \\citet{Jiang06} obtained the extinction around 7$\\mum$ for 129 different sightlines and no clear variation with Galactic longitude was found, although the extinction ratio $A_{[7]}/A_\\Ks$ does appear to exhibit a tendency of decreasing toward the Galactic center where $|l|<2^{\\circ}$\\citep{Jiang06}. The GLIMPSE Legacy Program surveyed the Galactic plane, with a large area coverage ($|l|\\leq65^{\\circ}$) and a detection limit of $\\simali$15.5--13.0$\\magni$ from 3.6 to 8.0$\\mum$ \\citep{GQA}. It provides an opportunity to explore the systematic variation of interstellar extinction in the IR with Galactic longitude.  In this work, we explore whether the mid-IR extinction varies among sightlines and how it varies in different interstellar environments based on the \\emph{Spitzer}/GLIMPSE database. In \\S\\ref{DATA}, the GLIMPSE data used in this work is briefly described. \\S\\ref{method} presents the method adopted to derive the extinction. In \\S\\ref{tracer} we discuss the selection of two different types of tracers (i.e. red giants and red clump giants). \\S\\ref{results} reports the resulting extinction ratios $A_{\\lambda}/A_\\Ks$ and the mean extinction from the total 131 GLIMPSE fields. Also discussed in \\S\\ref{results} are the comparison of the extinction derived here with previous studies performed by \\citet{Indebetouw05} and \\citet{Flaherty07}, and the longitudinal variation of the extinction ratios $A_{\\lambda}/A_\\Ks$ as well as its relation with the Galactic spiral arms and the distribution of interstellar dust. In \\S6 we summarize our major conclusions. ",
        "conclusions": "Using red giants and red clump giants as tracers, we have derived $A_\\lambda/A_\\Ks$, the extinction (relative to the 2MASS $\\Ks$ band) in the four IRAC bands, [3.6], [4.5], [5.8] and [8.0]$\\mum$ for 131 GLIPMSE fields along the Galactic plane within $|l|\\leq65^{\\rm o}$, based on the data obtained from the \\textit{Spitzer}/GLIPMSE Legacy Program and the 2MASS Survey project. The principal results of this paper are the following: \\begin{enumerate} \\item The mean extinction in the IRAC bands (normalized to the 2MASS $\\Ks$ band), $A_{[3.6]}/A_\\Ks\\approx0.63\\pm0.01$, $A_{[4.5]}/A_\\Ks\\approx0.57\\pm0.03$, $A_{[5.8]}/A_\\Ks\\approx0.49\\pm0.01$, and $A_{[8.0]}/A_\\Ks\\approx0.55\\pm0.03$, exhibits little variation with wavelength and lacks the minimum at $\\simali$7$\\mum$ predicted from the standard interstellar grain model for $R_V=3.1$. This is consistent within errors with previous observational determinations based on {\\it ISO} and {\\it Spitzer} data and with that predicted from the grain model for $R_V=5.5$ of \\citet{Weingartner01}. \\item The wavelength dependence of interstellar extinction in the mid-IR varies from one sightline to another, suggesting that there may not exist a ``universal'' IR extinction law.  \\item There exist systematic variations of extinction with Galactic longitude which appears to correlate with the locations of spiral arms and with the variation of the 240$\\mum$ dust emission. This can be understood in terms of larger grain sizes (arising from coagulational growth), enhanced dust concentration, and higher starlight intensities in the spiral arm regions. \\end{enumerate}"
    },
    "0910/0910.3989_arXiv.txt": {
        "abstract": "Extrasolar planets are not named and are referred to only by their assigned scientific designation. The reason given by the IAU to not name the planets is that it is considered impractical as planets are expected to be common. I advance some reasons as to why this logic is flawed, and suggest names for the 403 extrasolar planet candidates known as of Oct 2009. The names follow a scheme of association with the constellation that the host star pertains to, and therefore are mostly drawn from Roman-Greek mythology. Other mythologies may also be used given that a suitable association is established. ",
        "introduction": "Since the discovery of the first extrasolar planet, around the star 51 Pegasi (Mayor \\& Queloz 1995), over 400 planets surrounding other stars have been discovered. It is no exaggeration to say that, for astronomy, the year of 1995 has a historic resonance with 1781, 1846, and 1930. However, unlike Uranus, Neptune, and Pluto, the almost totality of these extrasolar planets are known by no other name than the scientifically dry designations given to them.  It is my intent to make the case that naming these planets is desirable. Poincar\\'e (1905) emphasized the usefulness of astronomy by saying that ``it is useful because it raises us above ourselves, because it is great, because it is beautiful''. Planet MOA-2008-BLG-310-L b, a sub-Saturnian mass planet recently detected in the Galactic Bulge with the technique of microlensing (Janczak et al. 2009), certainly inspires this feeling of transcendence Poincar\\'e describes. But its name hardly helps on conveying it.  One of the main reasons I consider for naming the extrasolar planets is the Copernican Principle itself. Our place in the cosmos is not special in any way, so there is no reason why only the planetary objects in the solar system should be named. Shakespeare would perhaps disagree with me and say that Io by any other name would smell as bad; and it is true that HD 128311 b will have the same radial velocity curve irrespectively of us naming it after a catalog number or after Bacchus. However, the non-special nature of our place in the Universe is better underscored by naming our neighbors. Mercury - Venus - Earth - Mars is a sequence of equals. Sol b - Sol c - Earth - Sol d would implicitly imply that the Earth is special in some way. Likewise, Jupiter is being paired to obscure names such as XO-1 b, TrES-4 b, and OGLE-TR-182 b, which does not help educators convey the message that these planets are quite similar to Jupiter. In stark contrast, the sentence``planet Apollo is a gas giant like Jupiter'' is heavily - yet invisibly - coated with Copernicanism.  One reason given by the IAU for not considering naming the extrasolar planets is that it is a task deemed impractical. One source is quoted as having said ``if planets are found to occur very frequently in the Universe, a system of individual names for planets might well rapidly be found equally impracticable as it is for stars, as planet discoveries progress.'' {\\footnote{\\url{http://www.iau.org/public_press/themes/extrasolar_planets/}}}. This leads to a second argument. It is indeed impractical to name all stars. But some stars are named nonetheless. In fact, all other classes of astronomical bodies are named. Galaxies are named. Stars are named. Even asteroids are named. So why not name the exoplanets? Granted, not all galaxies have names. Only the nearby ones, and in this case, the names are easily fixed due to the shape of the galaxy - Whirlpool, Antenna, Sombrero -, or of the constellation - Andromeda, Circinus, Carina. Star naming also has a criterion - the brightness. All stars brighter than $V$=1.5 have a proper name, the frequency of named stars declining as the magnitude increases{\\footnote{S. Boscardin, from Observat\\'orio Nacional, Rio de Janeiro, Brazil, pointed to me that there are four unnamed stars brighter than $V$=2.5. They are gamma and delta Velorum, of magnitudes $V$=1.8 and 2.0, and epsilon and eta Centauri, both of magnitude $V$=2.5. I credit him here for this information, and echo his suggestion that the IAU should consider naming these four lone stars. He further informs me that there are 84 stars between 2.5$<$$V$$<$3.0, 51 of which are named; and 106 stars between 3.0$<$$V$$<$3.5, 33 of which are named.}}. Yet nothing has stopped people from naming over 15,000 asteroids and minor planets. In fact, it seems to be that the main reason for over 400 known exoplanets remaining unnamed is that no one has yet done the job of naming them all. Indeed, as discoveries proceed (and hopefully skyrocket with the Gaia mission), naming all planets will be impractical. But the benefits of having some of them named is as clear as in the case of stars. In this manuscript, I set myself to the task.  In some cases, planets have already been nicknamed. For some, the epithet is sound. On naming 51 Pegasi b, I immediately thought of Bellerophon, the rider of the winged horse. Then I found out that someone else also shared the same taste for mythological associations and had already nicknamed the planet with the same name. I also came across the webpage of the ``Extrasolar planet naming society'', created in November 2008 with the idea of organizing a concerted collective effort to name the known exoplanets. Unfortunately, it did not seem to have gained any momentum. Individual efforts seem to beat collective ones in this case. I highlight here the webpage of Devon Moore, where over 40 exoplanet names have been suggested {\\footnote{\\url{http://nuclearvacuum.wikia.com/wiki/Names_for_extrasolar_planets}}}. Some of the names Moore chose are indeed good, some I have reasons to object. I will go back to that later.  Discoverers have also attempted to name the fruit of their labor. The three planets around Upsilon Andromedae have been nicknamed ``Fourpiter'', ``Twopiter'', and ``Dinky''. These names were suggested by a class of 4th graders, and have the advantage of carrying information on the planets' size.  However, they are tantalizingly unpalatable to those of a more classical mind. No offense to the very creative children, but apparently long gone are the times of Venitia Burney, then 11 years old, who suggested the name of Pluto, the Roman-Greek god of the underworld (not the Disney dog), for the distant, cold, (ex-)planet discovered in 1930. 70 Virginis b was similarly nicknamed ``Goldilocks'', for its position in the habitable zone of its star. One could imagine that had the discovery been made by a Swedish team if they would have called the planet {\\it Lagom}, after their unique word for ``just right''. Not to mention the internal names used by discovery teams that sometimes leak to the media, such as ``Xena'', ``Easterbunny'' and ``Santa'' (the Kuiper Belt objects Eris, Makemake, and Haumea, respectively).  In this manuscript, I advocate a return to classic tradition, and propose a simple way to name the known exoplanets, exposed in Sect 2. It basically consists of giving them names from Roman-Greek mythology associated with the constellation that the host star pertains. The columns in Table~1 show, respectively, the names of the planets listed by constellation, the name used in the astronomical literature, their masses in units of Jupiter's mass ($M_J$), semi-major axes in astronomical units (AU), eccentricities, right ascension, declination, and, finally, the proposed name{\\footnote{Table~1 can also be found online at \\\\\\url{http://www.mpia.de/homes/lyra/planet_naming.html}}}. The mythological associations for these names are explained in Sect 3.  The method used was the following. I got the planets from Jean Schneider's {\\it The Extrasolar Planets Encyclopaedia}, and listed the planets by constellation. I then suggest names for the planets with the help of a dictionary of Greek mythology (Dixon-Kennedy 1998) and extensive use of Wikipedia. It is commonplace that information in Wikipedia is not always trustable, and has to be used with care. That is correct, but when well-handled, the information there available is vital. For instance, listing the planets by constellation would have been quite time-consuming without Wikipedia, since the papers not always mention it, giving only the coordinates. To my surprise and amazement, almost all extrasolar planets have a wiki, that also informs in which constellation the star is located. I also made use of the following online source, \\url{http://www.theoi.com/}, that works as an online interactive dictionary of Greek mythology. ",
        "conclusions": "In this manuscript I suggest names for the 403 extrasolar planet candidates known as of October 2009. The suggestion is based on the classical tradition of giving names from Roman-Greek mythology to astronomical bodies. The association with the myths is in many cases purposively loose, to enable more flexibility on naming further planets as they are discovered. As said in \\sect{subsect:nonroman}, the system does not exclude other mythologies, which may be used if a suitable association with the constellation can be established.  The system also has some power of prediction. We can say for instance, that the next planet discovered in Eridanus could be called Minos, after one of the three judges of Hades. Indeed, in a former version of this manuscript I had suggested that the next three planets could be called Minos, Eachus, and Radhamantus. Then, on Sep 17, the Encyclopedia of Extrasolar Planets was updated. One of the new planets was in Eridanus, and I promptly added it to Table 1 as Radhamantus. In the update of Oct 9, another planet in Eridanus was announced. I added it to Table 1 as Eachus. Along the course of writing this manuscript, I could already experience another small benefit of having the planets named. Scrolling the long table up and down, it was much easier to remember ``Typhon'' than ``HD 168443 c''  I further stress that the system proposed here does not intend to supplant the one in vogue. It is not a change, but an addition. Stars are known by many names. Merope is also known as 23 Tau, HD 23480, HIC 17608, HR 1156, 2MASS J03461958+2356541, and V971 Tau, to name a few. The current naming scheme of assigning minor letters to the names of stars will of course be kept for scientific publications, much in the same way that we use HD 135742 instead of Zubenelschamali and NGC 4755 instead of Jewel Box. The proper name is a bonus aimed at popular writings.  One drawback I can think of is that it may lead the public to assume that the constellations are somehow physical. In some way, the misconception already exists. I recall this anecdote, for which I unfortunately cannot find the reference, that a piece of popular scientific writing once defined a constellation as ``a group of stars. Up to date, astronomers have found only 88''. On the other hand, we can invert the argument and see it as a golden opportunity to fight this misconception. Along with the names and their associations, it has to be pointed that the constellations are human invention and just a useful way of mapping the sky. Naming the planets after the myth of the constellations is no more misleading than using stellar names such as 14 Herculi, 70 Virginis or Upsilon Andromedae.  It should also be pointed that this manuscript was sent to the IAU Commission 53 on exoplanets, whose majority still opposes the idea of adding names to planets. The strongest concern of the commission had to do with the definition of an exoplanet. The Encyclopedia of Extrasolar Planets, complete as it is, lists candidate planets, and some of them are not confirmed. Therefore, some of the candidates assigned names here may as well be just low luminosity stars, brown dwarfs, a stellar spot or, as noted by a member of the commission, even a mote of dust in the spectrograph. In that case, the name should be withdrawn and re-assigned to other, confirmed, planet.  \\vspace{1cm}  {\\it Acknowledgments.} I thank the reading and comments of J. Alves, S. Boscardin, A. Johansen, M.-M. Mac Low, A. Moitinho, N. Piskunov, H. Rocha-Pinto, S. Soter, and P. Tsalmantza, that helped me on identifying the main points of critique. Interestingly, Portuguese speakers expressed concern regarding the use of the Lusiad. Westerns were concerned with the supposed Eurocentrism of mainly using Roman-Greek myths. It is great to see that we are living in multicultural times and that instead of trying to inflate national prides and highlighting boundaries, we are actually trying to distance ourselves from them. I also thank M. Schmitz for organizing an unorthodox peer review among the IAU Comission 53 and playing the role of editor.  \\vspace{1cm}"
    },
    "0910/0910.0278_arXiv.txt": {
        "abstract": "Non-radial pulsations (NRPs) are a proposed mechanism for the formation of decretion disks around Be stars and are important tools to study the internal structure of stars. NGC 3766 has an unusually large fraction of transient Be stars, so it is an excellent location to study the formation mechanism of Be star disks. High resolution spectroscopy can reveal line profile variations from NRPs, allowing measurements of both the degree, $l$, and azimuthal order, $m$. However, spectroscopic studies require large amounts of time with large telescopes to achieve the necessary high S/N and time domain coverage. On the other hand, multi-color photometry can be performed more easily with small telescopes to measure $l$ only. Here, we present representative light curves of Be stars and non-emitting B stars in NGC 3766 from the CTIO 0.9m telescope in an effort to study NRPs in this cluster. ",
        "introduction": "Be stars are a class of non-supergiant B-type stars with Balmer and other line emission features due to an equatorial decretion disk.  The disk is likely the result of a combination of the star's rapid rotation (near the critical limit) and non-radial pulsations (NRPs; Porter \\& Rivinius 2003).  NRPs are spherical harmonic waves traversing the surface of a star.  These pulsations can be found in multiple frequencies on the surface simultaneously (Rivinius, Baade, \\& \\u Stefl 2003).  There are two primary classes of NRP modes:  $g$- and $p$-modes.  $g$-modes are described by a low frequency pulsation that has gravity as its restoring force.  The dominant oscillation in this mode is transverse across the surface.  $p$-modes are dominated by high frequency, radial oscillations with a pressure restoring force (De Ridder 2001).  These modes in main-sequence, pulsating B stars are driven by the $\\kappa$ mechanism (Guti\\'errez-Soto et al. 2007).  Temperature and flux gradients are established between the dimmer, cooler material on the peaks of the pulsations and the brighter, warmer material in the troughs.  The flux variations over the stellar surface are then observed as either ripples within photospheric absorption line profiles or as periodic variations in magnitude.  A large, high-resolution spectroscopic study would reveal both the degree, $l$,  and the azimuthal order, $m$, but such studies are challenging due to the need for large amounts of time on a large telescope.  Photometry, which is easily performed with data gathered by small telescopes, only measures $l$ (Rivinius, Baade, \\& \\u Stefl 2003).  McSwain et al. (2008) previously showed that NGC 3766, an open cluster in Centaurus, is rich with transient Be stars.  In an effort to detect NRPs and study the formation of these transient disks, we are currently performing a long-term photometric study of the cluster.  Here we present preliminary differential light curves that reveal magnitude variations of several Be stars that are consistent with NRPs. ",
        "conclusions": ""
    },
    "0910/0910.2757_arXiv.txt": {
        "abstract": "We have performed C$^{18}$O ($J$=1--0) mapping observations of a $20'\\times20'$ area of the OMC-1 region in the Orion A cloud. We identified 65 C$^{18}$O cores, which have mean radius, velocity width in FWHM, and LTE mass of 0.18$\\pm$0.03 pc, 0.40$\\pm$0.15 km s$^{-1}$, and 7.2$\\pm$4.5 $M_\\odot$, respectively. All the cores are most likely to be gravitationally bound by considering the uncertainty in the C$^{18}$O abundance. We derived a C$^{18}$O core mass function, which shows a power-law-like behavior above 5 $M_\\odot$. The best-fit power-law index of $-2.3\\pm0.3$ is consistent with those of the dense core mass functions and the stellar initial mass function (IMF) previously derived in the OMC-1 region. This agreement strongly suggests that the power-law form of the IMF has been already determined at the density of $\\sim10^{3}$ cm$^{-3}$, traced by the C$^{18}$O ($J$=1--0) line. Consequently, we propose that the origin of the IMF should be searched in tenuous cloud structures with densities of less than 10$^{3}$ cm$^{-3}$. ",
        "introduction": "\\label{introduction} One of the most important observational features of the stellar initial mass function (IMF) is its power-law-like nature above 1 $M_\\odot$, as $dN/dM \\propto M^{-\\gamma}$. In the solar neighborhood, the power-law index $\\gamma$ seems to be greater than 2 \\citep{sal55,kro01a}, which characterizes the statistical properties of stars. In particular, both the total number and mass of stars are dominated by those of low-mass stars of $\\sim$ 1 $M_\\odot$. It is natural to consider that the origin of the IMF shape is related to the mass distribution of its natal gas in molecular clouds. Many authors have investigated dense gas ($10^{4-5}$ cm$^{-3}$) in molecular clouds by using (sub)millimeter dust continuum emission and/or molecular line emissions having high critical densities such as the H$^{13}$CO$^{+}$($J$=1--0) and N$_{2}$H$^{+}$($J$=1--0) lines \\citep[e.g.,][]{mot98,rei06,nut07,ike07,wal07,eno08}. They identified numerous cores that have typical sizes of 0.05 -- 0.1 pc and masses of 1 -- 10 $M_{\\odot}$ in nearby ($\\leq$ 500 pc) star forming regions such as Orion, Ophiuchus, Perseus, and Serpens. The molecular line studies showed that the cores are gravitationally bound and are likely to produce stars. Moreover, they found that the core mass functions (CMFs) derived by using the dense gas tracers, referred to as DCMFs hereafter, seem to have power-law-like behaviors in high-mass parts, whose $\\gamma$ values are very similar to that of the IMF. One exception is that \\citet{kra98} found a significatly smaller $\\gamma$ value of 1.7 for the S140 and M17SW regions by using the C$^{18}$O($J$=2--1) line with a high critical density comparable to that of H$^{13}$CO$^{+}$(1--0). Considering that the power-law form of the IMF has been already determined at the formation stage of the cores with the densities of $10^{4-5}$ cm$^{-3}$, it is likely that more tenuous structures of molecular clouds have a key to understanding the origin of the power-law nature of the IMF.  It has been suggested that the mass functions in the tenuous gas structures % are different from the DCMFs. \\citet{kra98} carried out a systematic study of the mass functions in the tenuous gas structures of 10$^{3}$ cm$^{-3}$ or less by using the $^{12}$CO(2--1), $^{13}$CO(1--0; 2--1), and C$^{18}$O(1--0) lines in various molecular clouds. They showed that their mass functions seem to have a common power-law form, and the $\\gamma$ value of 1.7$\\pm$0.1 is significantly smaller than those of the DCMFs and the IMF. \\citet{hei98} derived $^{12}$CO($J$=1--0, 2--1) mass functions in the MCLD 123.5$+$24.9 and Polaris Flare regions and found that $\\gamma$ = $1.8\\pm0.1$. \\citet{won08} found $\\gamma$ = 1.7 in the C$^{18}$O($J$=1--0) mass function of RCW 106. These facts mean that the power-law index of the IMF may be originated in the formation process of the dense gas of $10^{4-5}$ cm$^{-3}$ from the tenuous gas of $10^{2-3}$ cm$^{-3}$. However, one should be careful in comparing the $\\gamma$ value of the tenuous gas mass function to those of the DCMFs and the IMF. This is because the tenuous gas mass functions described above were derived by the spatial resolutions larger than 0.1 pc, which cannot resolve star-forming cores, and/or by using optically-thick tracers such as $^{12}$CO and $^{13}$CO. To fairly compare the tenuous gas mass function with the DCMFs and the IMFs, one should achieve higher spatial resolutions than 0.1 pc and use optically-thin tracers.  In this paper, we present a CMF derived by C$^{18}$O($J$=1--0) mapping observations of the OMC-1 region, which is located at the center of the Orion A Giant Molecular Cloud. The aim of this study is to examine whether or not the common power-law form between the DCMFs and the IMF has been already determined in the tenuous gas of $\\leq10^{3}$ cm$^{-3}$ by focusing on the power-law index $\\gamma$ in the high-mass part of the CMF. The C$^{18}$O($J$=1--0) molecular line emission is suitable for deriving the CMF \\citep[e.g.,][]{tac02} because the line has a relatively small critical density of $\\sim$10$^{3}$ cm$^{-3}$ \\citep{ung97} and is typically optically thin (see \\S \\ref{coreidentification}). The OMC-1 region is one of the best regions to investigate the CMF, because the IMF of the associated Orion Nebula Cluster (ONC) has been derived \\citep{hil97,mue02}, and because the power-law form of the H$^{13}$CO$^{+}$ DCMF is shown to be quite similar to that of the ONC IMF by \\citet{ike07}. Furthermore, at the distance to the Orion A cloud of 480 pc \\citep{gen81} we can easily resolve the cores with radii of $\\sim$0.1 pc. As described in \\S \\ref{observation}, the mapping observations have been done with the effective spatial resolution of 26$''$.4 ($=$ 0.06 pc), which is high enough to resolve the dense cores in the OMC-1 region \\citep{ike07}. ",
        "conclusions": ""
    },
    "0910/0910.0108_arXiv.txt": {
        "abstract": "The open clusters' fundamental physical parameters are important tools to understand the formation and evolution of the Galactic disk and as grounding tests for star formation and evolution models. However only a small fraction of the known open clusters in the Milky Way has precise determination of distance, reddening, age, metallicity, radial velocity and proper motion. One of the major problems in determining these parameters lies on the difficulty to separate cluster members from field stars and to assign membership. We propose a decontamination method by employing 2MASS data in the encircling region of the clusters NGC1981, NGC2516, NGC6494 and M11. We present a decontaminated CMD of these objects showing the membership probabilities and structural parameters as derived from King profile fitting. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.2611_arXiv.txt": {
        "abstract": "{Using numerical ray tracing, the paper studies how the average distance modulus in an inhomogeneous universe differs from its homogeneous counterpart. The averaging is over all directions from a fixed observer not over all possible observers (cosmic), thus is more directly applicable to our observations. In contrast to previous studies, the averaging is exact, non-perturbative, and includes all non-linear effects. The inhomogeneous universes are represented by Swiss-cheese models containing random and simple cubic lattices of mass-compensated voids. The Earth observer is in the homogeneous cheese which has an Einstein - de Sitter metric. For the first time, the averaging is widened to include the supernovas inside the voids by assuming the probability for supernova emission from any comoving volume is proportional to the rest mass in it. Voids aligned along a certain direction give rise to a distance modulus correction which increases with redshift and is caused by cumulative gravitational lensing. That correction is present even for small voids and depends on their density contrast, not on their radius. Averaging over all directions destroys the cumulative lensing correction even in a non-randomized simple cubic lattice of voids. At low redshifts, the average distance modulus correction does not vanish due to the peculiar velocities, despite the photon flux conservation argument. A formula for the maximal possible average correction as a function of redshift is derived and shown to be in excellent agreement with the numerical results. The formula applies to voids of any size that: (1) have approximately constant densities in their interior and walls; and (2) are not in a deep nonlinear regime. The average correction calculated in random and simple cubic void lattices is severely damped below the predicted maximal one after a single void diameter. That is traced to cancellations between the corrections from the fronts and backs of different voids. The results obtained allow one to readily predict the redshift above which the direction-averaged fluctuation in the Hubble diagram falls below a required precision and suggest a method to extract the background Hubble constant from low redshift data without the need to correct for peculiar velocities.} ",
        "introduction": "It is fair to say that the standard cosmological $\\Lambda$CDM model is facing a phenomenological crisis. The dark matter carrier has been evading direct detection for decades and the origin of dark energy remains a theoretical puzzle. The most natural candidate is the vacuum state energy, but the flat-space Quantum Field Theory is incapable of calculating it, producing an estimate that is $10^{120}$ times as large as the value suggested by the supernova observations. Presumably, that would be rectified in a fully quantized theory of gravitation which unfortunately does not yet exist. Remaining within the standard General Relativity, there are attempts to explain dark energy as an apparent quantity arising from the averaging procedure that maps the real lumpy universe to the idealized homogeneous Friedman-Lemaitre-Robertson-Walker (FLRW) metric. Such approaches separate into three distinct classes. First, one could average the Einstein equations over spatial hypersurfaces and hope the analogue of the Friedman equation contains a significant effective $\\Lambda$ term. Such an idea is conceptually problematic since inhomogeneous universes do not naturally pick up preferred spatial slices to average over, unlike the homogeneous FLRW models. It is shown in \\cite{wald} that the result of such averaging depends on the arbitrary choice of slicing and has no coordinate independent meaning - even Minkowsky spacetime could produce an apparent acceleration. A cousin to the first approach is averaging over the null hypersurface of the past light cone of the observer, \\cite{lightcone, valerioth}. It is not proven that such averages have a physical meaning and are not simply coordinate quantities. In the second approach, the second order perturbations to the Einstein equations are calculated and treated as an effective energy momentum tensor term. Unfortunately, as pointed out in \\cite{wald}, the second order perturbation term is not gauge independent and therefore is not suitable to play the role of an energy momentum tensor in the Einstein equation.  The present paper belongs to the third approach which is completely coordinate independent since it calculates relations only between physically observable quantities such as redshift and luminosity distance (or its logarithmic measure - the distance modulus).  Papers of this type, for example \\cite{norwegian},  demonstrate  that the distance modulus - redshift Hubble diagram of type Ia supernovas can be reproduced without any dark energy for an observer occupying the center of a giant ($\\sim$ Gpc) underdense spherical bubble. That is possible because the central observer measures a local Hubble parameter that is larger than the global one. The corresponding shift on the Hubble diagram between the true background and the one assumed by the observer allows for fitting the supernova data. Bubbles of such a large scale have significant peculiar velocities away from their center. That creates a fine tuning problem - the observer has to be in the vicinity of the bubble center \\cite{offcenter} to observe a small CMB dipole equal to the measured $v/c\\sim 0.002$ (Local Group velocity $\\approx620$ km/s). This violates the Copernicus principle that we do not occupy a special position in the universe.  Other papers of the same class studied the collective effect of a configuration of many bubbles/voids of smaller size, thus avoiding big peculiar velocities and significant anisotropies in CMB. That scenario can be conveniently simulated in the \"Swiss-cheese\" toy model \\cite{lemaitre, plebanski, bondi} which has the virtue of being analytically solvable. It is constructed by removing spherical regions from a homogeneous FLRW background (the \"cheese\") and replacing them with inhomogeneous density distributions with the same gravitating mass (mass-compensated voids). It was used in \\cite{valerio, valerioth} for an observer looking along the diameters of a string of aligned voids of radius $350$ Mpc. A positive cumulative change in the distance modulus was found which increased with redshift with respect to the homogeneous background. Although the size of the correction was significant at high redshift, it was not large enough to fit the supernova Hubble diagram and it substituted only partially for dark energy. A smaller correction due to aligned voids with a more mundane radius of $23 h^{-1}$ Mpc was calculated in \\cite{brouzakis-eqs}. The result obtained in \\cite{valerio} was criticized in \\cite{vanderveld} where it was demonstrated to stem from the cumulative weak lensing defocusing of the rays passing diametrically through the aligned voids. The same paper showed that the correction to the distance modulus vanishes when averaged over randomized impact parameters of the rays entering the voids. This conclusion is in accord with an old argument \\cite{weinberg} based on the gravitational lensing conservation of the total photon flux. The  implicit geometrical assumptions in \\cite{weinberg} have been challenged in more recent papers \\cite{mustapha, rose} (and references therein) and the present work will demonstrate they are violated close to the observer due to peculiar velocity modifications of the redshift surfaces.  The studies mentioned above are deficient in several areas. Averaging over directions was not performed in \\cite{valerio}. The calculations in \\cite{vanderveld} were done using weak lensing theory neglecting the time dependence of the void density and possible nonlinear effects which become significant at low redshifts. The averaging assumed that the ray impact parameters had a uniform probability distribution over the cross-section of a void. That is increasingly true at high redshifts but does not hold in the local neighborhood where the rays have to converge on the observer. More importantly, \\cite{vanderveld} averaged only over supernovas residing in the homogeneous cheese. The same was done in a study that obtained an analytical estimate of the corrections to the luminosity distance, \\cite{brouzakis}. Such a bias is unjustified, first because supernovas in the voids produce much bigger corrections than the ones in the cheese \\cite{valerio}, and second, the observed supernovas occur more frequently in denser regions like the void walls than in the cheese.  The goal of the present paper is to address once again the question whether the direction-averaged distance modulus correction from a configuration of voids really vanishes but this time in the context of a calculation that: (1) is exact, non-perturbative, an includes all possible non-linear effects; (2) is armed with a physically sensible averaging procedure free of ad hoc assumptions about impact parameters \\cite{vanderveld} or random cancellations of corrections \\cite{weinberg}; and (3) does not neglect the supernovas inside voids.  The final outcome of the calculations cannot be guessed on the grounds of the previous studies because it may depend on non-linear effects, the choice of averaging procedure, and the supernova selection, none of which was taken into account before. The calculations are performed in Swiss-cheese models with mass-compensated voids of two radii having observational support: $30$ and $300$ Mpc. The cheese is spatially-flat matter-only Einstein - de Sitter (EdS). The observer is placed in the cheese since at present there is no indication that we occupy a void. The first Swiss-cheese model is set up in section 2. The observer shoots past-directed light rays in all directions. The light propagation geodesic equations and the luminosity distance tracing along each ray are discussed in section 3. The physically sensible averaging procedure in this paper assumes that the probability for supernova emission from a comoving volume is proportional to the rest mass in it. Averaging the distance modulus for a single void is discussed in section 4 and it carries the essential characteristics of the procedure for many voids. The main outcome of that section is a simple formula to estimate the maximal average corrections to the distance modulus. Section 5 studies the cumulative correction due to gravitational lensing along a string of aligned voids. Numerical results from averaging within random and simple cubic void lattices are presented in section 6. The effect on the distance modulus produced by large voids of radius $300$ Mpc, which leave a measurable imprint on CMB, is evaluated in Section 7. The summary and conclusions section discusses the answer to the main question addressed by the present study and the practical importance of the obtained low-redshift results for future surveys.   The speed of light in this paper is $c=1$, and times and distances are measured in megaparsecs (Mpc), occasionally giving time in megayears (Myr) for convenience. The gravitational constant $G$ is kept explicit in all equations so that the reader can easily substitute with a favorite value. The usual geometrized units, $G=1$, were used in the numerical calculations. ",
        "conclusions": "The paper studied how the distance modulus, averaged over all lines of sight in an inhomogeneous universe, differs from its homogeneous counterpart. The inhomogeneities were represented by random and regular lattices of mass-compensated voids in Swiss-cheese models. Two void layouts were considered with a small ($30$ Mpc) and a big ($300$ Mpc) radius, the first observed in the SDSS data \\cite{visualvoid} and the second deduced from its imprint on CMB \\cite{granett}.  The Earth observer was put in the cheese but the conclusions will not change significantly for an observer inside one of the voids. For the first time, the averaging of the distance modulus was widened to include the supernovas inside voids. That was made possible by assuming that the probability of supernova emission from a comoving volume is proportional to the rest mass in it. Unlike previous studies \\cite{weinberg, vanderveld}, the average correction to the distance modulus was calculated by exact numerical ray tracing without using assumptions about the ray impact parameters or perturbation theory approximations, and the result includes all linear and non-linear effects.  The distance modulus correction, due to gravitational lensing, that accumulates along a ray crossing diametrically a sequence of aligned voids \\cite{valerio, brouzakis-eqs, vanderveld} was calculated. The correction increases linearly and becomes quite significant at high redshifts, see Fig.\\ref{cumucorr}. It is roughly proportional to the density contrast in the void interior but does not depend on the void radius. It was demonstrated that the distance modulus, averaged over all directions in a lattice of voids, does not show a cumulative correction at high redshifts. That is true not only in a random lattice (see Fig.\\ref{randvoidsave})  but also in a regular simple cubic one (see Fig.\\ref{scvoidsave} and Fig.\\ref{scvoidsave300}) implying that void randomization is not necessary for the destruction of the cumulative correction when averaging, contrary to popular belief. A large cumulative correction is still observed along special directions in a regular void lattice \\cite{valkenburg}, but the probability for the required void alignment is vanishing in the real universe.  At low redshifts $z$, perturbation theory predicts that the line-of-sight peculiar velocities $v_r$ are capable of producing significant fluctuations in the luminosity distance $\\Delta d_L(z)/d_L \\approx - v_r/z$ \\cite{hui, haugbolle} ($v_r$ measured in  speed of light units). That leads to a scatter in the local Hubble diagram $\\Delta H/ H = -\\Delta d_L/d_L$ and, by  (\\ref{fraccorr}), to a non-vanishing average correction $\\left<\\Delta \\mu(z)\\right>\\; \\approx \\, -2.17 \\left<\\Delta H / H \\right>   \\, \\approx \\, - 2.17 \\left<v_r(z)\\right>/\\,\\, z \\, $, where the so called peculiar velocity monopole $\\left<v_r(z)\\right>$ is the line-of-sight peculiar velocity at redshift $z$, averaged over all directions. A measure of the fluctuation magnitude is its average within a sphere of radius $R$ around the observer, $[\\Delta H/H]_R$,  which varies with location. Its cosmic variance over all possible observers/locations has been calculated previously in several cosmological models  by N-body simulations \\cite{edwin} and by linear perturbation theory using the matter power spectrum \\cite{shiturner, shidursi}. Such variances are employed to  predict confidence intervals for the value of $[\\Delta H/H]_R$ measured by a random observer. Since the obtained cosmic variance of $[\\Delta H/H]_R$ decreases with $R$ \\cite{shiturner, shidursi}, with the cosmic mean being zero, the likelihood to measure large average fluctuations decreases with redshift, as expected.  %  The cosmic variance of the sphere-average $[\\Delta H/H]_R$ over all observers cannot be used to estimate the average over all directions $\\left<\\Delta H(z)/H\\right>$ in the redshift bin $z$, experienced by a \\textit{particular} observer.  The results in the present paper allow to calculate $\\left<\\Delta H(z)/H\\right>$ by studying such an observer in a simple model of our local neighborhood. The averages and variances so obtained are not cosmic, thus are more applicable to our astronomical observations. They cannot be calculated from a given matter power spectrum but require a concrete numerical realization of an inhomogeneous universe. To the best knowledge of the author, this is the first time when the non-vanishing direction-averaged correction $\\left<\\Delta \\mu(z)\\right>$ is calculated as a function of redshift for universes containing voids. The upper bound (\\ref{corr}) for $\\left<\\Delta \\mu(z)\\right>$ was motivated using linear perturbation theory and was found to agree with the numerical calculations on Fig.\\ref{maxcorr} and Fig.\\ref{maxcorr300}.  The parameter $\\eta$ depends on the particular form of inhomogeneities and the choice of averaging procedure. The numerically obtained value $\\eta \\sim 0.20$ indicates that the maximal average correction is $20 \\%$ of the naive correction corresponding to the maximal peculiar velocity. This $\\eta$ turned out to be a universal constant, approximately independent of the void size, for voids which have approximately constant densities in their interior and walls that are not in a deep nonlinear regime. Due to void randomization, it is expected that the average correction will be suppressed below the maximal $20 \\%$ and will tend to zero. The efficiency of that process and at what redshift it sets in has not been studied before. The calculations revealed that the average correction within void lattices is indeed severely damped below the predicted maximal average correction due to cancelations between the fronts and backs of different voids. As figures \\ref{randvoidsave}, \\ref{scvoidsave}, and \\ref{scvoidsave300} demonstrate, the cancelation is surprisingly efficient at low redshifts even in regular lattices and the average correction drops below $0.01$ mag after a single void diameter. Nevertheless, the average correction is not zero close to the observer, indicating that the implicit assumptions of the photon flux conservation argument stated in \\cite{weinberg} do not apply at low redshifts.  The nonzero $\\left<\\Delta \\mu(z)\\right>$ is observable, provided there are enough supernovas $N_z$ in redshift bin $z$ to reduce the sampling statistical error of the average: $\\sigma[\\,\\left<\\Delta \\mu(z)\\right>\\,] = \\sigma[\\Delta\\mu(z)] /\\sqrt{N_z}$, where $\\sigma[\\Delta\\mu(z)]$ is shown as a thin solid curve on figures \\ref{randvoidsave}, \\ref{scvoidsave}, and \\ref{scvoidsave300}. An underdense \"Hubble bubble\" around us with a Hubble parameter slightly higher than the global one (correspondingly $\\left<\\Delta \\mu(z)\\right> < 0$) has been argued for in \\cite{zehavi, jha}. However, that claim was not confirmed in \\cite{giovanelli, neill} and was shown to depend on the way the supernova magnitudes were extracted from the photometric data \\cite{hicken}. Local bubble or not, fluctuations in the distance modulus binned average begin to appear in the newest data sets, Fig.20 in \\cite{kessler},  and Fig.7 in \\cite{sandage}.  Significant efforts  in contemporary cosmology are aimed at measuring a  possible time evolution in the dark energy equation of state. That  would require fixing the Hubble constant at low redshifts to a $1 \\%$ precision \\cite{jha, riess}. Ongoing and future low redshift surveys will boost the number of the observed nearby supernovas sufficiently to bring the sampling error of the average $\\sigma[\\left<\\Delta H / H \\right>]\\approx \\; \\sigma[\\,\\left<\\Delta \\mu(z)\\right>\\,] \\, / \\, 2.17 $ below $1 \\%$. In contrast, the directional average itself $\\left<\\Delta H / H \\right> \\approx \\;- \\left<\\Delta \\mu(z)\\right> \\, / \\, 2.17 $, generated by the coherent peculiar motion, is not influenced by the supernova number and will degrade significantly the errors in measuring the dark energy time dependence \\cite{hui, cooray}. Consequently, the peculiar velocities need to be subtracted from the low redshift Hubble diagram.  Several methods are utilized for that.  The Local Group (LG) of galaxies moves as a whole with respect to CMB at a speed $v_{LG}$, the so called peculiar velocity dipole \\cite{kocevski, watkins}. The most common way to correct the Hubble diagram for that motion is to adjust the observed redshifts to an observer riding the LG barycenter. The luminosity distance fluctuations seen in that frame are $\\Delta d_L/d_L \\approx -(v_r-v_{LG})/z$ \\cite{hui}  and the velocity dipole will cancel out since the coherent velocity component of a nearby  galaxy with respect to CMB is $v_r^{coh}\\approx v_{LG}$. Not correcting the CMB redshifts to the LG frame is permissible for a large supernova sample that covers densely and uniformly all directions: the coherent dipole fluctuation seen in the CMB frame is $\\Delta d_L/d_L \\approx- v_r/z\\approx -v_{LG}\\; cos(\\theta)\\, /z$ and that expression vanishes when averaged over the full solid angle. A further refinement is to correct the LG redshifts for infall velocities in the LG frame caused by the nearby superclusters represented by a crude linear multi-attractor model. That was done in the Hubble Key Project \\cite{mould} which measured the Hubble constant to a $9\\, \\%$ precision, but it was not very efficient at reducing the scatter in the low redshift Hubble diagram as seen on Fig.4 of \\cite{finalhubble}. Accordingly, the philosophy of the Hubble Key Project was to extract the Hubble constant mainly from secondary distance indicators at high redshifts. A finer method to correct for peculiar velocities is to calculate them from the observed galaxy distribution utilizing linear perturbation theory and choosing a galaxy-dark matter biasing parameter that minimizes the scatter on the Hubble diagram. That technique does reduce the scatter \\cite{radburn} but using it to infer the Hubble parameter from low redshift data produces a controversially high value of $H_0= 85 \\,km\\, s^{-1}\\,Mpc^{-1}$ \\cite{willick}. To avoid unreliable velocity corrections, many recent surveys \\cite{kessler, riess, sandagehubble} simply chose to include only objects above a certain redshift, $z>z_{min}$,  where the peculiar velocities are considered insignificant compared to the Hubble flow. The Hubble constant and the dark energy equation of state parameter $w$ inferred from the low redshift data are sensitive to the choice of $z_{min}$ \\cite{jha, kessler}. As discussed in \\cite{kessler}, there is a wide variation in the values selected for $z_{min}$ in the literature, ranging from $0.01$ \\cite{sandagehubble} up to $0.023$ \\cite{riess}, which indicates that there is not a universally accepted prescription.  Figures \\ref{randvoidsave}, \\ref{scvoidsave}, and \\ref{scvoidsave300}  of the present paper allow to select $z_{min}$ readily. The required $\\Delta H/H = 0.01$ precision corresponds to $\\Delta \\mu_{goal} = 0.0217$. The sampling error of the average is $\\sigma[\\Delta\\mu(z)] /\\sqrt{N_z}$, where $\\sigma[\\Delta\\mu(z)]$ is given by the thin solid curves on the plots and $N_z$ is the number of objects/supernovas in redshift bin $z$. Provided there is sufficient statistics, the sampling error will fall below $\\Delta \\mu_{goal}$. In that case the limiting value $z_{min}$ is determined by the average $\\left<\\Delta \\mu(z)\\right>\\ $ itself - a natural choice is the redshift at which the maximal average correction (\\ref{corr}) (the dashed curves) drops below $\\Delta \\mu_{goal}$ or the redshift corresponding to a void diameter (where $\\left<\\Delta \\mu(z)\\right>\\ \\sim 0.01 $ due to front-back void cancellation), whichever is lower. In comparison, estimates based on the sphere-average $[\\Delta H/H]_R$ are more pessimistic: its cosmic variance  in the $\\Lambda$CDM model drops below $1\\%$ at $z_{min}=0.05$ \\cite{shidursi} (linear estimate). If our neighborhood contains voids of the size observed on the CMB imprint \\cite{granett}, the required $z_{min}$ would be much larger, as Fig. \\ref{scvoidsave300} indicates.  Excluding the objects below $z_{min}$ reduces the size of the sample and proportionally increases the statistical error \\cite{cooray}. The present paper suggests it is possible to preserve the low redshift data, instead of throwing it away, by collapsing it into an interval average $[\\Delta \\mu]$ over a redshift interval containing a void diameter. These averages turn out very close to zero: $[\\Delta \\mu]=0.003$ ($z=0\\, \\ldots \\,0.025$ on Fig.\\ref{randvoidsave}), $[\\Delta \\mu]=0.007$ ($z=0\\, \\ldots \\, 0.025$ on Fig.\\ref{scvoidsave}), $[\\Delta \\mu]=0.002$ ($z=0\\, \\ldots \\,0.25$ on Fig.\\ref{scvoidsave300}). That stems from the cancellation between positive and negative lobes in $\\left<\\Delta \\mu(z)\\right>$ and the mass averaging giving higher weights to higher redshifts where $\\left<\\Delta \\mu(z)\\right>$ is closer to zero. A  vanishing $[\\Delta \\mu] \\approx 0$ means $[\\mu]\\approx[\\mu]^{EdS}$:  \\begin{equation} \\label{extracthubble} \\sum_i W_i \\left<\\mu(z_i)\\right> \\approx \\sum_i W_i \\; \\mu_i^{EdS}(H_0, z_i), \\end{equation}  where $W_i$ is the mass weight assigned to redshift bin $z_i$, $\\left<\\mu(z_i)\\right>$ is the distance modulus in that bin averaged over all directions, and the sum is over all the bins in the redshift interval. The lefthand side of the formula is calculated from the observational data. The weight $W_i$ is proportional to the rest mass inside the redshift bin which can be estimated as the corresponding mass in the homogeneous background model. Unlike a traditional Hubble diagram that weighs all redshift bins equally, here $W_i \\propto z_i^2 \\, \\Delta z_i$ at low redshifts. The righthand side of (\\ref{extracthubble}) depends solely on the Hubble parameter $H_0$ at low redshifts and on additional cosmological parameters of the background model at higher redshifts. It allows one to extract the value of $H_0$ from the low redshift data. This method will become possible for future surveys that: (1) have a dense full-sky coverage to calculate the average correction over all directions $\\left<\\Delta \\mu(z)\\right>$; and (2) contain a large number of objects in each redshift bin so that $\\left<\\Delta \\mu(z)\\right>$ is readily observable, not swamped by the sampling error.  If the observer is inside a void, the average correction $\\left<\\Delta \\mu(z)\\right>$ starts with a negative lobe unlike the so-far considered case of an outside observer. A simple illustration of that is an observer at the center of a void, which will measure $\\Delta \\mu(z) = \\mu(z) - \\mu^{EdS}(z)<0$ for supernovas inside the void since the matter there is expanding faster than the background and it takes a smaller distance $d_L$ to achieve the same redshift. The interval average $[\\Delta \\mu]$ over a void diameter would still be very small since the mass averaging scheme adopted in this paper ascribes lower weights to low redshifts. Additional studies are needed to investigate the behavior of $[\\Delta \\mu]$ in the presence of superclusters - the other type of major inhomogeneities encountered in the universe. A plausible guess is that $[\\Delta \\mu]$ vanishes on a redshift interval encompassing a supercluster, justifying the validity of (\\ref{extracthubble}) in that case as well."
    },
    "0910/0910.1612_arXiv.txt": {
        "abstract": "{We study the morphological content of a large sample of high-redshift clusters to determine its dependence on cluster mass and redshift. Quantitative morphologies are based on PSF-convolved, 2D bulge+disk decompositions of cluster and field galaxies on deep Very Large Telescope FORS2 images of eighteen, optically-selected galaxy clusters at $0.45 < z < 0.80$ observed as part of the ESO Distant Cluster Survey (``EDisCS''). Morphological content is characterized by the early-type galaxy fraction $f_{et}$, and early-type galaxies are objectively selected based on their bulge fraction and image smoothness. This quantitative selection is equivalent to selecting galaxies visually classified as E {\\it or} S0. Changes in early-type fractions as a function of cluster velocity dispersion, redshift and star-formation activity are studied. A set of 158 clusters extracted from the Sloan Digital Sky Survey is analyzed exactly as the distant EDisCS sample to provide a robust local comparison. We also compare our results to a set of clusters from the Millennium Simulation. Our main results are: (1) The early-type fractions of the SDSS and EDisCS clusters exhibit no clear trend as a function of  cluster velocity dispersion. (2) Mid-$z$ EDisCS clusters around $\\sigma$ = 500 km/s have $f_{et} \\simeq$ 0.5 whereas high-$z$ EDisCS clusters have $f_{et} \\simeq$ 0.4. This represents a $\\sim$25$\\%$ increase over a time interval of 2 Gyrs. (3) There is a marked difference in the morphological content of EDisCS and SDSS clusters. None of the EDisCS clusters have early-type galaxy fractions greater than 0.6 whereas half of the SDSS clusters lie above this value.  This difference is seen in clusters of all velocity dispersions. (4) There is a strong and clear correlation between morphology and star formation activity in SDSS and EDisCS clusters in the sense that decreasing fractions of [OII] emitters are tracked by increasing early-type fractions.  This correlation holds independent of cluster velocity dispersion and redshift even though the fraction of [OII] emitters decreases from $z \\sim0.8$ to $z \\sim 0.06$ in all environments. Our results pose an interesting challenge to structural transformation and star formation quenching processes that strongly depend on the global cluster environment (e.g., a dense ICM) and suggest that cluster membership may be of lesser importance than other variables in determining galaxy properties. ",
        "introduction": "Our current paradigm for the origin of galaxy morphologies rests upon hierarchical mass assembly \\citep[e.g.,][]{steinmetz02}, and many transformational processes are at work throughout the evolutionary histories of galaxies. Some determine the main structural traits (e.g., disk versus spheroid) while others only influence properties such as color and star-formation rates. Disk galaxy collisions lead to the formation of elliptical galaxies \\citep{spitzer51,toomre72,farouki82,negroponte83,barnes92,barnes96,mihos96}, and the extreme example of this process is the build-up of the most massive galaxies in the Universe at the cores of galaxy clusters through the accretion of cluster members. Disks can also be transformed into spheroidals by tidal shocks as they are harassed by the cluster gravitational potential \\citep{farouki81,moore96,moore98}. Harassment inflicts more damage to low luminosity galaxies because of their slowly rising rotation curves and their low density cores. Galaxies can be stripped of their internal gas and external supply through ram pressure exerted by the intracluster medium \\citep{gunn72,larson80,quilis00}, and the result is a ``quenching'' (or ``strangulation'') of their star formation that leads to a rapid reddening of their colours \\citep[also see][]{martig09}. The task of isolating observationally the effects of a given process has remained a major challenge to this day.  Many processes affecting galaxy morphologies are clearly environmentally-driven, and galaxy clusters are therefore ideal laboratories in which to study all of them. The dynamical state of a cluster, which can be observationally characterized by measuring mass and substructures, should be related to its morphological content. For example, the number of interactions/collisions suffered by a given galaxy should depend on local number density and the time it has spent within the cluster. Dynamically young clusters with a high degree of subclustering should contain large numbers of galaxies that are infalling for the first time. More massive clusters will contain more galaxies, but they will also have higher galaxy-galaxy relative velocities that may impede merging \\citep{lubin02}. Spheroidal/elliptical galaxies will preferentially be formed in environments where the balance between number density and velocity dispersions is optimal, but it is still not clear where this optimal balance lies. Cluster masses can be estimated from their galaxy internal velocity dispersion \\citep{rood72,dressler84,carlberg97,tran99,borgani99,lubin02}, through weak-lensing shear \\citep{kaiser93,schneider95,hoekstra00,clowe06} or through analysis of their hot X-ray emitting atmospheres \\citep[e.g., ][]{allen98}, and it will be used here as the main independent variable against which morphological content will be studied.  The morphological content of high-redshift clusters is most often characterized by the fraction $f_{E+S0}$ of early-type galaxies they contain \\citep{dressler97,dokkum00, fasano00, dokkum01, lubin02, holden04, smith05, postman05, desai07,poggianti09b}. The bulk of the data available so far is based on visual classification. ``Early-type'' galaxies are defined in terms of visual classifications as galaxies with E or S0 Hubble types. A compilation of early-type fractions taken from the literature \\citep{dokkum00} shows a dramatic increase of the early-type fractions as a function of decreasing redshift from values around 0.4$-$0.5 at $z \\sim 1$ to values around 0.8 in the local Universe. However, the interpretation of this trend is not entirely clear as others \\citep[e.g.,][]{dressler97,fasano00,desai07,poggianti09b} have reported that the fraction of E's remains unchanged as a function of redshift and that the observed changes in early-type fractions are entirely due to the S0 cluster populations. S0 populations were observed to grow at the expense of the spiral population \\citep{smith05,postman05,moran07,poggianti09b} although others \\citep[e.g.,][]{holden09} have argued for no evolution in the relative fraction of ellipticals and S0s with redshift. \\citet{smith05} and \\citet{postman05} show that the evolution of $f_{E+S0}$ is in fact a function of both lookback time (redshift) and projected galaxy density. They find $f_{E+S0}$ stays constant at 0.4 over the range 1 $< t_{lookback} <$ 8 Gyr for projected galaxy densities $\\Sigma <$ 10 Mpc$^{-2}$. For high density environments ($\\Sigma$ = 1000 Mpc$^{-2}$), $f_{E+S0}$ decreases from 0.9 to 0.7. At fixed lookback time, $f_{E+S0}$ varies by a factor of 1.8 from low to high densities at $t_{lookback}$ = 8 Gyr and by a factor of 2.3 at $t_{lookback}$ = 1 Gyr. The difference between low and high density environments thus increases with decreasing lookback time. Both studies indicate that the transition between low and high densities occurs at $0.6R_{200}$ ($R_{200}$ is the projected radius delimiting a sphere with interior mean density 200 times the critical density at the cluster redshift, see Equation~\\ref{radius200}). \\citet{postman05} also find that $f_{E+S0}$ does not change with cluster velocity dispersion for massive clusters ($\\sigma$ $>$ 800 km/s). The data for one of their clusters also suggest that $f_{E+S0}$ decreases for lower mass systems.  This trend would be consistent with observations of $f_{E+S0}$ in groups that show a strong trend of decreasing $f_{E+S0}$ versus decreasing $\\sigma$ \\citep{zabludoff98}. Finally, $f_{E+S0}$ seems to correlate with cluster X-ray luminosity at the 2-3$\\sigma$ level \\citep{postman05}.  Recent works on stellar mass-selected cluster galaxy samples \\citep{holden07,vanderwel07} paint a different picture. The fractions of E+S0 galaxies in clusters, groups and the field do not appear to have changed significantly from $z \\sim 0.8$ to $z \\sim 0.03$ for galaxies with masses greater than 4$\\times 10^{10} M_{\\odot}$. The mass-selected early-type fraction remains around 90\\% in dense environments ($\\Sigma >$ 500 gal Mpc$^{-2}$) and 45\\% in groups and the field. These results show that the morphology-density relation of galaxies more massive than 0.5M$_{*}$ has changed little since $z \\sim 0.8$ and that the trend in morphological evolution seen in luminosity-selected samples must be due to lower mass galaxies. This is in agreement with \\citet{delucia04,delucia07} and \\citet{rudnick09} who have shown the importance of lower mass (i.e., fainter) galaxies to the evolution of the color-magnitude relation and of the luminosity function versus redshift. Another interesting result has come from attempts to disentangle age, morphology and environment in the Abell 901/902 supercluster \\citep{wolf07,lane07}. Local environment appears to be more important to galaxy morphology than global cluster properties, and while the expected morphology-density and age-morphology relations have been observed, there is no evidence for a morphology-density relation at a fixed age. The time since infall within the cluster environment and not density might thus be the more fundamental parameter dictating the morphology of cluster galaxies.  A number of efforts have been made on the theoretical side to model the morphological content of clusters. \\citet{diaferio01} used a model in which the morphologies of cluster galaxies are solely determined by their merger histories. A merger between two similar mass galaxies produces a bulge, and a new disk may form through the subsequent cooling of gas. Bulge-dominated galaxies are in fact formed by mergers in smaller groups that are later accreted by clusters. Based on their",
        "conclusions": "\\label{discussion}  In order to fully understand possible evolutionary trends observed here, it is important to determine how cluster velocity dispersion changes with redshift as a result of the hierarchical growth of structures. Are we looking at similar clusters when we focus on the same range of velocity dispersions in the SDSS and EDisCS clusters? \\citet{poggianti06} looked at the mean change in $\\sigma$ between $z$ = 0 and $z$ = 0.76 using a sample of 90 haloes from the Millennium Simulation uniformly distributed in log(mass) between 5 $\\times$ 10$^{12}$ and 5 $\\times$ 10$^{15}$ M$_{\\odot}$. Their Figure 8 shows how $\\sigma$ evolves over that redshift interval. For example, a $z$ = 0 cluster with $\\sigma$ = 900 km/s would typically have $\\sigma \\sim$ 750 km/s at $z$ = 0.76. This evolution is not sufficient to introduce biases in our analysis here. Indeed, selecting clusters with $\\sigma \\geq$ 600 km/s, say, at either $z$ = 0 or $z$ = 0.76 would keep nearly all the same clusters. Measured velocity dispersions may exhibit a large scatter with respect to the true halo mass particularly for low-mass clusters. The velocity dispersions for the SDSS and EDisCS clusters were calculated in a very similar way in order to minimize any biases. Velocity dispersions calculated from a small number of cluster members may be overestimates of the true cluster mass. Table~\\ref{clsample} lists 1103.7-1245b as the cluster with the lowest number of members ($N$ = 11). In order to check the robustness of our results, we re-ran our analyses by excluding SDSS clusters in Table~\\ref{sdss-cls-list} with $N < 10$ for which velocity dispersions may be less reliable and found that our results remained unchanged.  \\citet{poggianti06} proposed a scenario in which two channels are responsible for the production of passive galaxies in clusters, and others \\citep{faber07,brown07} have proposed a similar scenario for the migration of galaxies from the \"blue cloud\" to the red sequence. \"Primordial passive galaxies\" are composed of galaxies whose stars all formed at very high redshift ($z >$ 2) over a short timescale. These galaxies have been observed in clusters up and beyond $z = 1$, and they largely comprise luminous ellipticals. \"Quenched passive galaxies\" have had a more extended period of star formation activity, and their star formation has been quenched after their infall into dense cluster environments. These quenched passive galaxies would then suffer the effects of cluster processes such as ram pressure stripping, harassment, strangulation and mergers to become S0 and earlier type galaxies. A key point of this scenario is that processes affecting morphology and star formation activity operate on different timescales as shown recently for the EDisCS sample by \\citet{sanchez09}. There is good evidence that star formation is quenched in galaxies over timescales of 1-3 Gyr after they have entered the cluster environment \\citep{poggianti99,poggianti06} whereas morphological transformation through mergers and harassment can take longer \\citep[$\\sim$ 5 Gyr,][]{moore98}. The best example of this is the fact that the vast majority of post-starburst galaxies in distant clusters, those that have had their star formation activity terminated during the last Gyr, still retain a spiral morphology \\citep{poggianti99}. Such a two-channel scenario would naturally explain observations indicating that the elliptical galaxy fraction actually remains constant with redshift while the S0 fraction rises with decreasing redshift \\citep{dressler97,fasano00,desai07}. Unfortunately, the VLT/FORS2 images do not have sufficient spatial resolution to disentangle E and S0 galaxies as mentioned in Section~\\ref{efrac-defn} to determine the exact contribution from  each channel. We can therefore only study the overall production of early-type galaxies, but it should exhibit different behaviors with cluster global properties depending on the process(es) dominating it. Given our quantitative definition of an early-type galaxy based on bulge fraction and image smoothness, there are essentially two ways to transform late-type galaxies into early-type ones: 1) processes such as collisions and harassment that can fundamentally alter the structure of a galaxy by forming bulges and/or destroying disks and 2) quenching processes that can extinguish star forming regions responsible for some of the galaxy image asymmetries and also cause a fading of the disks.  Applying the \\citet{poggianti06} scenario to our results,  the \"threshold\" in $f_{et}$ values in our high redshift clusters (Figures~\\ref{vlt+sdss_efrac_sigma} and~\\ref{vlt+sdss_efrac_age}) could be explained by a population of primordial passive galaxies that formed at even higher redshifts. Most of our high redshift clusters have early-type fractions in the range 0.3-0.6 with no correlation with cluster velocity dispersion. Are these early-type fractions indeed consistent with a populations of primordial passive galaxies? Calculations done in \\citet{poggianti06} show that the fraction of galaxies at $z = 0.6$ that were present in haloes with masses greater than 3$\\times$10$^{12}$ M$_\\odot$ at $z$ = 2.5 is 0.4$\\pm$0.2. These primordial passive galaxies can therefore account for at least 2/3 (if not all) of the early-type populations in high redshift clusters, and their high formation redshift would explain the lack of dependence of $f_{et}$ on cluster velocity dispersion.  One of our main results is that the early-type fractions of galaxy clusters increase from $z = 0.6 - 0.8$ to $z\\sim 0.08$ in clusters of all velocity dispersions. What kind of morphological transformation process(es) can lead to such an evolution? Collisions and harassment both depend on galaxy-galaxy interactions and the time a galaxy has spent within the cluster environment. Cluster velocity dispersion influences the number of interactions and their duration. Higher velocity dispersions in more massive clusters yield more interactions per unit time $N$ but with shorter durations $\\Delta t$ in a given time interval. One might therefore expect to see a peak in early-type type fraction at the cluster velocity dispersion where the product N$\\Delta t$ is maximized. No such peak is seen in our clusters. Ram-pressure stripping is expected to go as ($n_{ICM} v_{gal}^{2.4})/\\dot{M}_{rep}$ \\citep{gaetz87} with $n_{ICM}$, $v_{gal}$ and $\\dot{M}_{rep}$ being the density of the ICM, the velocity of the galaxies within the ICM and the rate at which galaxies can replenish their gas respectively. The fraction of passive galaxies should therefore be a relatively strong function of cluster velocity dispersion if quenching by ram pressure stripping is the dominant process. The number of post-starburst galaxies in EDisCS clusters does correlate with cluster velocity dispersion \\citep{poggianti09a}, but the uniform increase in early-type fractions at all cluster velocity dispersions observed going from EDisCS to SDSS clusters is not consistent with the intracluster medium being the main cause of the changes in cluster morphological content.  Even though the early-type and [OII] emitter fractions in EDisCS and SDSS clusters show no correlation with cluster velocity dispersion \\citep[][and this work]{poggianti06}, there is a very strong correlation between $f_{et}$ and $f_{OII}$. This correlation is seen at both low and high cluster masses as well as at both low and high redshifts. Morphology and star formation therefore appear to be closely linked with one another over a wide range of environments and times. However, different structural transformation and quenching processes are thought to operate over different timescales \\citep[e.g.,][]{sanchez09}. Timescales range from 1-2 Gyr (based on typical cluster crossing times) for truncating star formation to 3-5 Gyr for totally extinguishing star formation in newly accreted galaxies \\citep{poggianti06,tonnesen09}. Looking at the evolution of EDisCS cluster red-sequence galaxies over 2 Gyr (from $z = 0.75$ to $z$ = 0.45), \\citet{sanchez09} found that morphological transformation and quenching of star formation indeed appeared to not be simultaneous. As noted in Section~\\ref{efrac-sigma}, the early-type fractions of mid-$z$ EDisCS clusters may be $\\sim$25$\\%$ higher than the ones of high-$z$ clusters. This change would therefore have taken place over a 2 Gyr interval in our adopted cosmology.  However, the time baseline here between SDSS and EDisCS clusters is almost 6 Gyr, and, unfortunately, this is ample time to erase any difference arising from different timescales in the link between morphology and star formation.  The lack of dependence of morphology and star formation on global cluster properties such as velocity dispersion raises the question of whether changes in galaxy properties are driven by more local effects or whether they occur outside of the cluster environment. Recent work \\citep{poggianti08,park09,bamford09,ellison09} have re-emphasized the strong link between galaxy properties and local galaxy density rather than cluster membership. Galaxy properties are seen to change at densities around 15-40 galaxies Mpc$^{-2}$ or projected separations of 20-30$h^{-1}$ kpc. Others \\citep[e.g.,][]{kautsch08,wilman09} have suggested that the galaxy group environment might be more conducive to galaxy transformation. Our observed evolution in early-type fraction as a function of redshift and the strong correlation between morphology and star formation at all cluster masses would support the idea that cluster membership is of lesser importance than other variables such as local density in determining galaxy properties.  The properties of simulated clusters from the Millenium Simulation compare well with those of EDisCS and SDSS clusters. Their early-type fractions also show no dependence with cluster velocity dispersion in contrast to previous theoretical work \\citep[e.g.][]{diaferio01} but in agreement with observations. However, there is a definite lack of MS clusters with low early-type fractions at $z$ = 0 compared to the SDSS sample. It is important here to note that an early-type galaxy in the simulations was defined solely based on its bulge fraction because the simulations do not have the resolution required to model internal fine structures such as asymmetries. Given that real, early-type galaxies were also selected according to image smoothness, one would expect the early-type fractions of real clusters to be systematically lower. However, half of the SDSS clusters have low early-type fractions not seen in the simulations at $z$ = 0, and such a large discrepancy could only be explained by a significant population of real bulge-dominated galaxies with relatively large asymmetries. It is more likely that bulge formation in the simulations may be too efficient. The scatter in $f_{et}$ values for the simulated clusters with $\\sigma \\geq $ 600 km/s is also nearly three times smaller than observed in the real clusters (Section~\\ref{efrac-sigma}) which may indicate that the models may not include the right mixture of evolutionary processes at work on real galaxies. High-mass simulated clusters show a correlation between early-type fraction and star-forming fraction (albeit over narrower ranges than observed), but the correlation is not seen in the low-mass simulated clusters. This may be understood by high mass clusters having been formed long enough for evolutionary processes to have had enough time to act on galaxies to modify their properties whereas this is not necessarily the case for low-mass clusters. The fact that the correlation is observed in both low- and high-mass real clusters may be an indication that processes giving rise to the correlation may be more efficient (or altogether different) than modelled. It is also important to keep in mind here that the properties of a galaxy in these models are essentially driven by the mass of its parent halo."
    },
    "0910/0910.4037_arXiv.txt": {
        "abstract": "{The seismic  data obtained by CoRoT for the star HD~49933 enable us for the first time to measure \\emph{directly}  the amplitudes and linewidths of solar-like oscillations for a star other than the Sun. From those measurements it is possible, as was done for the Sun, to constrain models of the excitation of acoustic modes by turbulent convection. } {We compare a stochastic excitation model described in Paper~I with the asteroseismology data for HD~49933, a star that is rather metal poor and significantly hotter than the Sun.} {Using the seismic determinations of the mode linewidths detected by CoRoT for HD~49933 and the theoretical mode excitation rates computed in Paper~I for the specific case of HD~49933,  we derive the expected surface velocity amplitudes of the acoustic modes detected in HD~49933. Using a calibrated quasi-adiabatic approximation relating the mode amplitudes in intensity to those in velocity, we derive the expected values of the mode amplitude in intensity.  } { Except at rather high frequency, our amplitude calculations are within 1-$\\sigma$  error bars of the mode surface velocity spectrum derived with the HARPS  spectrograph.  The same is found with respect to the mode amplitudes in intensity derived for HD~49933 from the CoRoT data. On the other hand, at high frequency ($\\nu \\gtrsim~1.9 $~mHz),  our calculations depart significantly  from the CoRoT and HARPS measurements. We show that assuming a solar metal abundance rather than the actual metal abundance of the star would result in a larger discrepancy with the seismic data. Furthermore, we present calculations which assume the ``new'' solar chemical mixture to be in better agreement with the seismic data than those that assumed the ``old''  solar chemical mixture.} {  These results validate in the case of a star significantly hotter than the Sun and $\\alpha$~ Cen~A the main assumptions in the model of stochastic excitation. However, the discrepancies seen at high frequency highlight some deficiencies of the modelling, whose origin remains to be understood. We also show that it is important to take  the surface metal abundance of the solar-like pulsators into account. } ",
        "introduction": "The amplitudes of solar-like oscillations result from a balance between excitation and damping. The mode linewidths are directly related to the mode damping rates. Once we can measure the mode linewidths, we can derive the theoretical value of the mode amplitudes  from theoretical calculations of the mode excitation rates, which in turn can be compared to the available seismic constraints. This comparison allows us to test the model of stochastic mode excitation investigated in a companion paper \\citep[][hereafter Paper~I]{Samadi09a}.  As shown in Paper~I, a moderate deficit of the surface metal abundance results in a significant decrease of the mode driving by turbulent convection. Indeed, by taking into account the measured iron-to-hydrogen abundance ([Fe/H]) of HD~49993 ([Fe/H]$=-0.37$), we have derived the theoretical values of the mode excitation rates ${\\cal P}$ expected for this star. The resulting value of ${\\cal P}$ is found to be about two times smaller than for a model with the same gravity and effective temperature, but with a solar metal abundance (i.e. [Fe/H]$=0$).  The star HD~49933 was first observed in Doppler velocity by \\citet{Mosser05} with the HARPS spectrograph. More recently, this star has been observed  twice by CoRoT. A first time this was done  continuously during about 61 days (initial run, IR) and a second time continuously during about 137 days (first long run in the center direction, LRc01). The combined seismic analysis of these data \\citep{Benomar09b} has provided the mode linewidths as well as the amplitudes of the modes in intensity. Then, using mode linewidths  obtained for  HD 49933 with the CoRoT data  and the theoretical mode excitation rates (obtained in Paper~I), we derive the expected values of the mode surface \\emph{velocity} amplitudes. We next compare these values  with the mode velocity spectrum derived following \\citet{Kjeldsen05} with seismic data  from the  HARPS spectrograph \\citep{Mosser05}.  Mode amplitudes in terms of \\emph{luminosity} fluctuations  have also been derived from the CoRoT data for 17 radial orders.  These data provide us with not only a constraint on the maximum of the mode amplitude but also with the frequency dependence. % The relative luminosity amplitudes $\\delta L/L$ are linearly related to the velocity amplitudes. This ratio is determined by the solution of the \\emph{non-adiabatic} pulsation equations and is independent of the stochastic excitation model \\citep[see][]{Houdek99}. Such a non-adiabatic calculation requires us to take into account, not only the radiative damping, but also the coupling between the pulsation and the turbulent convection.  However, there are currently very significant uncertainties concerning the modeling of this coupling \\citep[for a recent review see ][]{Houdek08}. We relate further for the sake of simplicity the mode luminosity amplitudes to computed mode velocity amplitudes by assuming adiabatic oscillations as \\citet{Kjeldsen95}. Such a relation is calibrated in order to reproduce the helioseismic data.   The comparison between theoretical values of the mode amplitudes (both in terms of surface velocity and intensity) constitutes a test of the stochastic excitation model  with a star significantly different from the Sun and {\\acenA}. In addition  it is also possible to test the validity of the calibrated quasi-adiabatic relation, since both mode amplitudes, in terms of surface velocity and intensity, are available for this star,  This paper is organized as follows: We describe in Sect.~\\ref{velocity} the way mode amplitudes in terms of surface velocity $v_s$ are  derived from the theoretical  values of ${\\cal P}$ and from the measured mode linewidths ($\\Gamma$). Then, we compare the theoretical values of the mode surface velocity with the seismic constraint obtained from HARPS observations. We describe in Sect.~\\ref{intensity} the way mode amplitudes in terms of intensity fluctuations $\\delta L/L$ are derived from theoretical  values of $v_s$ and compare  $\\delta L/L$ with the seismic constraints obtained from the CoRoT observations. Finally, Sects. \\ref{Discussion} and \\ref{Conclusion} are dedicated to a discussion and conclusion respectively. ",
        "conclusions": "\\label{Conclusion}  From the mode linewidths measured by CoRoT  and theoretical mode excitation rates derived for HD~49933, we have derived the expected mode surface velocities $v_s$ which we have compared with $v_{\\rm HARPS}$, the mode velocity spectrum derived from the seismic observations obtained with the HARPS spectrograph \\citep{Mosser05}. Except at high frequency ($\\nu \\gtrsim$~1.9~mHz), the agreement between  computed $v_s$ and $v_{\\rm  HARPS}$ is within the 1-$\\sigma$ domain associated with the seismic data from the HARPS spectrograph. However,  there is a clear tendency to overestimate $v_{\\rm  HARPS}$ above $\\nu~\\sim$~1.9~mHz,.  Using a \\emph{calibrated} quasi-adiabatic approximation to relate the mode velocity to the mode amplitude in intensity (Eq.~\\ref{dL_Vad}), we have derived for the case of HD~49933 the expected mode amplitudes in intensity. Computed mode intensity fluctuations, $\\delta L/L$, are within 1-$\\sigma$  in agreement with the seismic constraints derived from the CoRoT data \\citep{Benomar09b}. However, as for the velocity, there is a clear tendency at high frequency ($\\nu \\gtrsim$~1.9~mHz) towards over-estimated $\\delta L/L$ compared to the CoRoT observations.     Calculations that assume a solar surface metal abundance result, both in velocity and in intensity, in amplitudes larger by $\\sim$~35\\,\\% around the peak frequency ($\\nu_{\\rm max} \\simeq$ 1.8~mHz) and by up to a factor of two at lower frequency. It follows that, ignoring the current surface metal abundance of the star results in a more severe over-estimation of the computed amplitudes compared with observations. This illustrates the importance of taking the surface metal abundance of the solar-like pulsators into account when modeling the mode driving. In addition, we point out that the \\citet{GN93} solar chemical mixture results in mode amplitudes larger by about 15\\,\\% with respect to calculations that assume the ''new'' solar abundance by \\citet{Asplund05}. However, this increase remains significantly smaller than the  uncertainties associated with current seismic measurements.  Since both mode amplitudes in terms of surface velocity and intensity are available for this star, it was possible to test the validity of the calibrated quasi-adiabatic relation (\\eq{dL_Vad}). Our comparison  shows that this relation   provides the correct scaling, at least at the level of the present seismic precisions,  .  Both in terms of surface velocity and of intensity, the differences between predicted and observed mode amplitudes  are within the 1-$\\sigma$ uncertainty domain, except at high frequency. This result then validates for low frequency modes the basic underlying physical assumptions included in the theoretical model of stochastic excitation  for a star significantly different in effective temperature, surface gravity, turbulent Mach number ($M_t$) and metallicity compared to the Sun or $\\alpha$~Cen~A.  As discussed in Sect.~\\ref{Discussion}, the clear discrepancy  between predicted and observed mode amplitudes  seen at high frequency may have two possible origins: First, a canceling between the entropy contribution and the Reynolds stress is expected to occur and to be important around and above the frequency of the maximum of the mode excitation rates (see Sect.~\\ref{canceling_entropy_reynolds}). Second, the assumption called the ``scale length separation'' \\citep{Samadi08b} may also result in an over-estimation of the mode amplitudes at high frequency (see Sect.~\\ref{scale length separation}). These issues will be investigated in a forthcoming paper."
    },
    "0910/0910.3421_arXiv.txt": {
        "abstract": "\\vsco\\ is an eclipsing dwarf nova that had attracted little attention from X-ray astronomers until it was proposed as the identification of an \\xte\\ all-sky slew survey (XSS) source.  Here we report on the pointed X-ray observations of this object using \\suzaku.  \\vsco\\ was in quiescence at the time, as indicated by the coordinated optical photometry we obtained at the South African Astronomical Observatory.   Our \\suzaku\\ data show \\vsco\\ to be X-ray bright, with a highly absorbed spectrum. Most importantly, we have discovered a partial X-ray eclipse in \\vsco. This is the first time that a partial eclipse is seen in X-ray light curves of a dwarf nova.  Our preliminary simulations demonstrate that the partial X-ray eclipse can be in principle reproduced if the white dwarf in \\vsco\\ is partially eclipsed. Higher quality observations of this object have the potential to place significant constraints on the latitudinal extent of the X-ray emission region and thereby discriminating between an equatorial boundary layer and a spherical corona. The partial X-ray eclipse therefore makes \\vsco\\ a key object in understanding the physics of accretion in quiescent dwarf nova. ",
        "introduction": "Cataclysmic variables (CVs), in which a white dwarf primary accretes from a Roche-lobe filling, late-type secondary (see \\citealt{tome} for a comprehensive review), are an excellent laboratory for the physics of accretion.  In the subclass of dwarf novae,  accretion proceeds via a disk that switches between the low (quiescence) and high (outburst) luminosity states.  The visible light of a quiescent dwarf nova is usually dominated by the bright spot, where the accretion flow from the secondary hits the outer edge of the disk, as well as the photosphere of the white dwarf.  In outburst, the disk becomes the dominant source of visible light.  In contrast, X-ray observations of dwarf novae probe the accretion flow in the immediate vicinity of the white dwarf, such as an optically thin boundary layer \\citep{PR1985a}, since the accretion disk in a dwarf nova is too cool to emit X-rays.  Although the disk instability model is a highly successful framework for explaining the dwarf nova outbursts, there are details that defy the prediction of the basic version of the model \\citep{L2001}.  In particular, the observed X-ray luminosities of quiescent dwarf nova (often of order 10$^{31}$ \\eps; \\citealt{Bea2005}) imply an accretion rate onto the white dwarf that is much higher than predicted.   Proposed modifications of the disk instability model include the coronal siphon flow \\citep{MM1994}, which might lead to accretion over a much wider area of the white dwarf surface than through a boundary layer, and a weakly magnetic white dwarf \\citep{LP1992}.  In the latter model, the magnetic field is too weak to control the accretion flow during outburst, but strong enough to do so during quiescence.  \\vsco\\ is an eclipsing \\citep{BSG2000} dwarf nova with an orbital period of 1.8 hr \\citep{T1999}.  According to \\cite{KMU2002}, \\vsco\\ has a quiescent magnitude of $\\sim$14.5, and has an outburst every $\\sim$30 days during which it reaches magnitude $\\sim$12.5.   Extensive photometry by Warner and collaborators \\citep{WWP2003,PWW2006} have revealed quasi-periodic oscillations and dwarf nova oscillations in \\vsco, but a strictly periodic signal was never found.  Therefore, existing optical data point strongly towards a dwarf nova classification and argue against an intermediate polar (IP, or DQ Her type systems; \\citealt{P1994}) classification.  IPs are a subset of magnetic CVs in which the primary's magnetic field disrupts the inner accretion disk, channeling the flow to the magnetic polar region(s); the spin period of the magnetic white dwarf is a strict clock that characterize IPs. Thus, \\vsco\\ joins a growing number of eclipsing dwarf novae below the period gap.  Citing the variable shape of the eclipses and the relatively small eclipse amplitude (often less than 0.75 mag), \\cite{BSG2000} argue for a grazing eclipse of the bright spot and the disk, but not of the white dwarf.  On the other hand, \\cite{Mea2000} argue that the white dwarf is eclipsed, based on the fact that the spectroscopic conjunction of the disk (the red-to-blue crossing of the emission line radial velocities) occurs at mid-eclipse. From the radial velocity curves, they also infer a white dwarf mass of $\\sim$0.5--0.6 M$_\\odot$ and a mass ratio of $\\sim$0.2--0.3.  However, \\cite{Mea2001} prefer a higher mass (0.89 M$_\\odot$) white dwarf and an inclination angle of 72.5 degrees based on their analysis of the emission line radial velocities, although this value depends in part on the assumed mass-radius relationship for the secondary.  Despite the current lack of a consensus regarding the system parameters, the eclipsing nature of \\vsco\\ makes it an important target for detailed studies.  Moreover, it is a nearby system with parallax-estimated distance of 155$^{+58}_{-34}$ pc \\citep{T2003}.  Although the \\rosat\\ detection was already noted by \\cite{Kea1998}, \\vsco\\ did not draw the attention of X-ray astronomers until it was listed among the \\xte\\ all-sky slew survey (XSS) sources \\citep{Rea2004}, along with three other dwarf novae (SS~Aur, V426~Oph, and SU~UMa).  The estimated luminosities of these 4 systems are all just under 10$^{32}$ \\eps\\ in the 2--10 keV band, placing them near the upper end for non-magnetic CVs ($10^{30}$--$3 \\times 10^{32}$ in 0.1--100 keV for the \\asca\\ sample; \\citealt{Bea2005}). However, the XSS is based on the data taken during slews of the non-imaging \\xte\\ PCA detector.  For relatively faint sources such as these dwarf novae, the positional errors are considerable, and misidentification is a possibility.  Hence it is important to confirm the proposed identification of XSS dwarf novae.  This is particularly true for \\vsco, which had never been the subject of a pointed X-ray observation above 2 keV. Although \\vsco\\ is securely detected as a \\rosat\\ all-sky survey (RASS) source, the ratio of XSS (2.38 c/s in the 3--8 keV band) to RASS (0.35) count rates is high, compared to, e.g., SS~Aur which has the same RASS count rate but was detected at 0.75 c/s in the XSS 3--8 keV band. This leaves open the possibility that \\vsco\\ is only a partial contributor to the XSS flux.  For this reason and also because of the potential return in studying the X-ray properties of this eclipsing dwarf nova, we performed a pointed observation using \\suzaku, along with contemporaneous optical photometry. ",
        "conclusions": "We performed \\suzaku\\ observations of the eclipsing dwarf nova, \\vsco.  We confirm that it is an X-ray luminous dwarf nova, and moreover, report the discovery of a partial X-ray eclipse.  In the past, both optical and X-ray observers have concentrated on systems that exhibit a total eclipse of the white dwarf.  They provide a strong constraint on the height of the X-ray emission region above the white dwarf. However, eclipse light curves are one-dimensional, and one can only derive one-dimensional constraints from the eclipse analysis. The X-ray emission region, on the other hand, has two dimensions (height and latitudinal extent) assuming azimuthal symmetry. It is therefore necessary to observe an ensemble of eclipsing dwarf novae, in which the limb of the secondary cuts across the emission region from different angles, to be able to constrain the height and the latitudinal extent simultaneously.  Inclusion of a partially eclipsing system in such a study will be a huge step forward in our quest to constrain the X-ray emission region size in two dimensions.  Future X-ray observations of \\vsco, together with improved estimates of system parameters from other wavelengths, have the potential to test the boundary layer picture of X-ray emission in quiescent dwarf novae, and furthermore to constrain the latitudinal extent of such a boundary layer."
    },
    "0910/0910.4806_arXiv.txt": {
        "abstract": "We present a careful and detailed light curve analysis of RR Lyrae stars in the  Small Magellanic Cloud (SMC) discovered by the Optical Gravitational Lensing Experiment (OGLE) project. Out of  536 single mode RR Lyrae stars selected from the database, we have investigated the physical properties of 335  `normal looking' RRab stars and 17 RRc stars that have good quality photometric light curves. We have also been able to estimate the distance modulus of the cloud which is in good agreement with those determined from other independent methods. The Fourier  decomposition method has  been used to study the basic properties of these variables. Accurate Fourier decomposition parameters of 536 RR Lyrae stars in the OGLE-II database are computed. Empirical relations between the Fourier parameters and some physical parameters of these variables have been used to estimate the physical parameters for the stars from the Fourier analysis. Further, the Fourier decomposition of the light curves of the SMC RR Lyrae stars yields their mean physical parameters as: [Fe/H] = -1.56\\,$\\pm\\,0.25$, M = 0.55 $\\pm $\\,0.01\\,M$_{\\odot}$, T$_{\\rm eff} = 6404\\, \\pm 12$ K, $\\log \\rm L = 1.60\\,\\pm0.01\\,\\rm L_\\odot$ and M$\\rm _V = 0.78\\,\\pm0.02 $ for 335 RRab variables and [Fe/H] = -1.90\\,$\\pm$\\, 0.13, M = 0.82 $\\pm $\\,0.18\\,M$_{\\odot}$, T$_{\\rm eff} = 7177\\,\\pm 16$ K, $\\log\\,\\rm L = 1.62\\,\\pm \\,0.02\\,\\rm L_{\\odot}$ and M$\\rm _V = 0.76\\,\\pm 0.05$ for 17 RRc stars. Using the absolute magnitude together with the mean magnitude, intensity-weighted mean magnitude and the phase-weighted mean magnitude of the RR Lyrae stars, the mean distance modulus to the SMC is estimated to be 18.86\\,$\\pm$0.01 mag, 18.83\\,$\\pm$0.01 mag and 18.84\\,$\\pm$0.01 mag respectively from the RRab stars. From the RRc stars, the corresponding distance modulus values are found to be 18.92\\,$\\pm$0.04 mag, 18.89\\,$\\pm$0.04 mag and 18.89\\,$\\pm$0.04 mag respectively. Since Fourier analysis is a very powerful tool for the study of  the physical properties of the RR Lyrae stars, we emphasize the importance of exploring the reliability of the calculation of Fourier parameters together with the uncertainty estimates keeping in view the large collections of photometric light curves that will become available from variable star projects of the future. ",
        "introduction": "RR Lyrae (RRL) stars have played an important role in determining the cosmic distance scale in modern astronomy. While several studies have used classical Cepheids as primary distance indicators, RRL stars have received substantial attention in solving the distance scale problem.  Apart from distance determinations, RRL stars are of particular importance as a test bed for the theories of stellar and galactic structure and evolution. They are the tracers of old stellar populations of the bulge, disk and halo components that are present everywhere.  RRL stars are radially pulsating A-F variable stars with periods in the 0.2 - 1.2 day range, and amplitude of variation $ \\leq 2$ mag. They can be easily identified, and play a key role as the cornerstone of the Population II distance scale. They are extensively used to determine distances to old and sufficiently metal-poor systems, where they are commonly found in large numbers. In particular, RRL stars are present in globular clusters (GCs) and the dwarf galaxies in the neighborhood of the Milky Way (Greco et al. 2007), and have also been identified in the M31 field (Brown et al. 2004, Dolphin et al. 2004), in some M31 companions (Pritzl et al. 2005), and in  at least four M31 GCs (Clementini et al. 2001). Distances to  the Large Magellanic Cloud (LMC) for the population II objects are based on the luminosity of the RRL stars (Clementini et al. 2002).  There have been a very few studies to estimate the distance scale of the SMC using the RRL stars. Using four RRL light curves of the SMC cluster NGC 121, Walker \\& Mack  (1988) have estimated the distance modulus of the cluster to be 18.86$\\pm$0.07 mag. On the  other hand, using 22 RR Lyrae stars surveyed around 1.3 square degrees near the northeast arm of the SMC field NGC 361, Smith et al. (1992) obtained the distance modulus of the SMC as 18.90$\\pm$0.16 mag. Also, there are other studies of estimating the distance scale of the SMC based on the binary star light curves and double mode Cepheids. Harries et al. (2003) reported that the distance modulus to the SMC is of the order of 18.89$\\pm$0.14 mag taking 10 eclipsing binaries in the SMC. By selecting 40 eclipsing binaries of spectral type O and B in the SMC, Hilditch et al. (2005) have derived the fundamental parameters of the binaries and refined the distance modulus to the SMC to 18.91$\\pm$0.1 mag. Also Kov\\'{a}cs et al. (2000) found the distance modulus of the SMC to be 19.05$\\pm$0.017 mag based on the photometric data of double-mode Cepheids from the OGLE project.  The stellar atmospheric parameters of effective temperature (T$_{\\rm eff}$) and surface gravity (log g) are of fundamental astrophysical importance. They are the prerequisites to any detailed abundance analysis and define the physical conditions in the stellar atmosphere and hence are directly related to the physical properties mass (M), radius (R) and luminosity (L) of the star. In this paper, we present an independent analysis of 536 SMC RRL stars discovered by the OGLE project (Soszy\\'{n}sky et al. 2002). The OGLE database is a very wealthy resource for studying the characteristics of variable stars in the Galaxy, LMC and SMC. For the first time, we make use of the OGLE SMC data to estimate the distance scale of the SMC using a large number of well-sampled RRL light curves.  The road map of the present investigation is to perform a Fourier analysis of the RRL stars in order to estimate their physical parameters and hence the SMC distance scale.  We employ the Fourier decomposition technique which is used extensively to characterize the observed photometric light curves of RRL and other types of variables. The accurate determination of the Fourier coefficients is, therefore, an important task. We have performed  an independent automated Fourier analysis of all the RRL light curves selected in this paper by a computer code developed by us. In section 2 we give a brief description of the database that we use and the procedure of removing the outliers from the light curve data. We present Fourier decomposition of the light curves in section 3. We also describe the use of  the unit-lag auto-correlation function for finding out the optimal order of the fit to the RRL light curves. Section 4 describes the error analysis of the Fourier decomposition parameters $\\rm \\phi_{i1}$ and R$_{i1}$.  Section 5 describes the calibration of the I band  data to the  V band. In section 6, we describe the various physical parameters of the RRLs obtained by using empirical relations from the literature. Section 7 describes the distance determination of the SMC. Lastly, in section 8, we present the conclusions of our study. ",
        "conclusions": "In this paper we have derived the physical parameters of 352 RR Lyrae stars (335 RRab and 17 RRc) of the SMC from OGLE-II I band database using the Fourier decomposition of their light curves. The stars were selected based on the quality of their light curves. I band Fourier coefficients have been converted to V band using the inter-relations obtained from the observational data of B06 and W94. Using the Fourier decomposition method, we find the mean physical parameters: [Fe/H] = -1.56\\,$\\pm\\,0.25$, M = 0.55 $\\pm $\\,0.01\\,M$_{\\odot}$, T$_{\\rm eff} = 6404\\, \\pm 12$ K, $\\log \\rm L = 1.60\\,\\pm0.01\\,\\rm L_\\odot$ and M$\\rm _V = 0.78\\,\\pm0.02 $ for 335 RRab variables and [Fe/H] = -1.90\\,$\\pm$\\, 0.13, M = 0.82 $\\pm $\\,0.18\\,M$_{\\odot}$, T$_{\\rm eff} = 7177\\,\\pm 16$ K, $\\log\\,\\rm L = 1.62\\,\\pm \\,0.02\\,\\rm L_{\\odot}$ and M$\\rm _V = 0.76\\,\\pm 0.05$ for 17 RRc stars. Mean distance modulus to the SMC was calculated from the 352 light curves by using mean magnitude, intensity weighted mean magnitude and phase weighted mean magnitude. The values of the distance modulus are found to be in good agreement with independent studies.  Locations of the RRL stars in the H-R diagram clearly show that the estimates of the parameters determined from the Fourier decomposition method are consistent with the theoretical blue and red edges of the instability strip calculated by Bono et al. (1995). The calculations of the radii of the RRLs by using the Fourier decomposition technique are in agreement with the theoretical period-radius-metallicity relations of Marconi et al. (2005) obtained from a completely different approach of non-linear convective modelling. \\begin{figure} \\centering \\includegraphics[height=8cm,width=9cm]{fig17.eps} \\vspace{10pt} \\caption{Histogram of the metallicity of the RRLs used for the analysis. The partially filled histogram shows the metallicity distribution of the 335 RRab stars, whereas the fully filled histogram shows the metallicity distribution of the 17 RRc stars.} \\label{Fig 14}% \\end{figure}  \\begin{table*} \\caption{Physical parameters extracted from Fourier coefficients for 335 RRab variables. Errors represent the uncertainties in the Fourier parameters.} \\begin{center} \\scalebox{0.9}{ \\begin{tabular}{ccccccccccccccccc} \\\\ \\hline \\hline OGLE ID &M$_{V}$&\\rm $\\log$\\,(L/L$_{\\odot})$&[\\rm Fe/H]&T$_{\\rm eff}$&\\rm M/M$_{\\odot}$&\\rm log{\\rm  g}&$\\rm R/R_{\\odot} (FD)$& $\\rm R/R_{\\odot} $ (Mar)  \\\\ ~~~~~~~(1)&(2)&(3)&(4)&(5)&(6)&(7)&(8)&(9)\\\\  \\hline \\hline 004801.59-733021.5&   0.93$\\pm$   0.02&    1.53$\\pm$   0.01&   -0.91$\\pm$   0.19&  6751$\\pm$ 9&    0.61$\\pm$   0.01&    2.97$\\pm$   0.02&    4.27$\\pm$   0.02&    4.31$\\pm$   0.03\\\\ 005300.26-725136.6&   0.92$\\pm$   0.02&    1.54$\\pm$   0.01&   -1.18$\\pm$   0.19&  6702$\\pm$ 9&    0.64$\\pm$   0.01&    2.97$\\pm$   0.02&    4.40$\\pm$   0.02&    4.43$\\pm$   0.03\\\\ 003841.21-734422.9&   0.90$\\pm$   0.01&    1.54$\\pm$   0.01&   -1.00$\\pm$   0.18&  6758$\\pm$ 9&    0.60$\\pm$   0.01&    2.95$\\pm$   0.02&    4.33$\\pm$   0.02&    4.41$\\pm$   0.03\\\\ 003727.78-731454.9&   0.92$\\pm$   0.01&    1.53$\\pm$   0.01&   -0.88$\\pm$   0.16&  6755$\\pm$ 8&    0.59$\\pm$   0.01&    2.95$\\pm$   0.01&    4.29$\\pm$   0.02&    4.38$\\pm$   0.02\\\\ 005728.85-723454.6&   0.93$\\pm$   0.02&    1.52$\\pm$   0.01&   -0.76$\\pm$   0.21&  6739$\\pm$10&    0.57$\\pm$   0.01&    2.94$\\pm$   0.02&    4.26$\\pm$   0.02&    4.36$\\pm$   0.03\\\\ 005728.85-723454.6&   0.93$\\pm$   0.02&    1.52$\\pm$   0.01&   -0.76$\\pm$   0.21&  6739$\\pm$10&    0.57$\\pm$   0.01&    2.94$\\pm$   0.02&    4.26$\\pm$   0.02&    4.36$\\pm$   0.03\\\\ 005026.32-732418.2&   0.94$\\pm$   0.02&    1.52$\\pm$   0.01&   -1.00$\\pm$   0.20&  6653$\\pm$ 9&    0.60$\\pm$   0.01&    2.94$\\pm$   0.02&    4.39$\\pm$   0.02&    4.48$\\pm$   0.03\\\\ 004639.18-731324.7&   0.89$\\pm$   0.02&    1.54$\\pm$   0.01&   -0.81$\\pm$   0.20&  6800$\\pm$10&    0.56$\\pm$   0.01&    2.93$\\pm$   0.02&    4.29$\\pm$   0.02&    4.42$\\pm$   0.03\\\\ 003816.46-732449.2&   0.92$\\pm$   0.02&    1.52$\\pm$   0.01&   -0.74$\\pm$   0.19&  6738$\\pm$ 9&    0.55$\\pm$   0.01&    2.93$\\pm$   0.02&    4.29$\\pm$   0.02&    4.42$\\pm$   0.03\\\\ 005110.48-730750.0&   0.87$\\pm$   0.01&    1.55$\\pm$   0.01&   -0.85$\\pm$   0.18&  6805$\\pm$ 9&    0.55$\\pm$   0.01&    2.92$\\pm$   0.02&    4.33$\\pm$   0.02&    4.48$\\pm$   0.03\\\\ 010452.90-724025.9&   0.95$\\pm$   0.01&    1.52$\\pm$   0.01&   -1.00$\\pm$   0.19&  6599$\\pm$ 9&    0.59$\\pm$   0.01&    2.92$\\pm$   0.02&    4.44$\\pm$   0.02&    4.55$\\pm$   0.03\\\\ 005646.16-723452.2&   0.88$\\pm$   0.01&    1.55$\\pm$   0.01&   -1.06$\\pm$   0.18&  6691$\\pm$ 9&    0.57$\\pm$   0.01&    2.90$\\pm$   0.02&    4.48$\\pm$   0.02&    4.65$\\pm$   0.03\\\\ 005957.83-730647.6&   0.87$\\pm$   0.01&    1.55$\\pm$   0.01&   -0.99$\\pm$   0.17&  6717$\\pm$ 8&    0.56$\\pm$   0.01&    2.90$\\pm$   0.01&    4.44$\\pm$   0.02&    4.63$\\pm$   0.03\\\\ 005458.09-724948.9&   0.89$\\pm$   0.01&    1.55$\\pm$   0.00&   -1.11$\\pm$   0.15&  6646$\\pm$ 7&    0.58$\\pm$   0.01&    2.90$\\pm$   0.01&    4.51$\\pm$   0.01&    4.68$\\pm$   0.02\\\\ 004758.98-732241.3&   0.89$\\pm$   0.02&    1.54$\\pm$   0.01&   -0.95$\\pm$   0.26&  6683$\\pm$12&    0.56$\\pm$   0.01&    2.90$\\pm$   0.02&    4.45$\\pm$   0.03&    4.63$\\pm$   0.04\\\\ 010535.93-720621.6&   0.94$\\pm$   0.02&    1.53$\\pm$   0.01&   -0.98$\\pm$   0.30&  6583$\\pm$13&    0.57$\\pm$   0.01&    2.89$\\pm$   0.03&    4.49$\\pm$   0.03&    4.67$\\pm$   0.05\\\\ 010516.55-722526.5&   0.90$\\pm$   0.02&    1.53$\\pm$   0.01&   -0.76$\\pm$   0.20&  6702$\\pm$ 9&    0.53$\\pm$   0.01&    2.89$\\pm$   0.02&    4.38$\\pm$   0.02&    4.59$\\pm$   0.03\\\\ 004721.26-731135.5&   0.89$\\pm$   0.01&    1.55$\\pm$   0.01&   -1.02$\\pm$   0.17&  6640$\\pm$ 8&    0.56$\\pm$   0.01&    2.88$\\pm$   0.02&    4.52$\\pm$   0.02&    4.73$\\pm$   0.03\\\\ 005504.67-731106.4&   0.91$\\pm$   0.01&    1.54$\\pm$   0.01&   -1.08$\\pm$   0.19&  6583$\\pm$ 9&    0.57$\\pm$   0.01&    2.88$\\pm$   0.02&    4.56$\\pm$   0.02&    4.76$\\pm$   0.03\\\\ 004306.71-733527.9&   0.92$\\pm$   0.02&    1.53$\\pm$   0.01&   -0.89$\\pm$   0.20&  6605$\\pm$10&    0.54$\\pm$   0.01&    2.88$\\pm$   0.02&    4.49$\\pm$   0.02&    4.70$\\pm$   0.03\\\\   \\hline \\hline \\end{tabular} } \\end{center} Complete table is available in the electronic form. \\end{table*}    \\begin{table*} \\caption{Physical parameters extracted from Fourier coefficients for 17 RRc variables. Errors represent the uncertainties in the Fourier parameters.  } \\scalebox{0.9}{ \\begin{tabular}{cccccccccc} \\\\ \\hline \\hline OGLE ID &\\rm M$_{\\rm V}$(K98)&$\\rm \\log (\\rm L/L_{\\odot}) (SC93) $&[\\rm Fe/H]&T$_{\\rm eff}$&\\rm M/M$_{\\odot}$&\\rm log \\rm g&$\\rm R/R_{\\odot} (FD)$&$\\rm R/R_{\\odot}$ (Mar)  \\\\ ~~~~~~~(1)&(2)&(3)&(4)&(5)&(6)&(7)&(8)&(9) \\\\ \\hline \\hline 005451.72-723850.4&   0.85$\\pm$   0.03&    1.66$\\pm$   0.02&   -1.10$\\pm$   0.17& 7457$\\pm$17&    0.59$\\pm$   0.19&    3.00$\\pm$   0.33&    4.07$\\pm$   0.03&    4.19$\\pm$   0.02\\\\ 005115.64-724739.2&   0.78$\\pm$   0.04&    1.67$\\pm$   0.02&   -1.21$\\pm$   0.16& 7425$\\pm$17&    0.62$\\pm$   0.19&    3.00$\\pm$   0.31&    4.19$\\pm$   0.04&    4.25$\\pm$   0.02\\\\ 010245.99-721132.7&   0.83$\\pm$   0.04&    1.83$\\pm$   0.03&   -2.02$\\pm$   0.11& 7165$\\pm$18&    1.13$\\pm$   0.07&    3.04$\\pm$   0.06&    5.35$\\pm$   0.05&    4.63$\\pm$   0.02\\\\ 004631.02-724658.8&   0.78$\\pm$   0.04&    1.75$\\pm$   0.03&   -1.67$\\pm$   0.17& 7266$\\pm$20&    0.77$\\pm$   0.18&    2.97$\\pm$   0.24&    4.79$\\pm$   0.05&    4.65$\\pm$   0.02\\\\ 004902.13-724513.6&   0.82$\\pm$   0.03&    1.76$\\pm$   0.02&   -1.73$\\pm$   0.13& 7247$\\pm$15&    0.80$\\pm$   0.18&    2.97$\\pm$   0.23&    4.87$\\pm$   0.04&    4.69$\\pm$   0.02\\\\ 010453.59-723756.0&   0.78$\\pm$   0.03&    1.83$\\pm$   0.01&   -2.11$\\pm$   0.08& 7115$\\pm$10&    0.93$\\pm$   0.17&    2.94$\\pm$   0.18&    5.44$\\pm$   0.04&    5.09$\\pm$   0.01\\\\ 010029.57-725454.2&   0.77$\\pm$   0.04&    1.81$\\pm$   0.02&   -2.03$\\pm$   0.09& 7136$\\pm$10&    0.85$\\pm$   0.19&    2.93$\\pm$   0.22&    5.32$\\pm$   0.05&    5.10$\\pm$   0.01\\\\ 010029.57-725454.2&   0.77$\\pm$   0.04&    1.81$\\pm$   0.02&   -2.03$\\pm$   0.09& 7136$\\pm$10&    0.85$\\pm$   0.19&    2.93$\\pm$   0.22&    5.32$\\pm$   0.05&    5.10$\\pm$   0.01\\\\ 004327.25-724542.4&   0.76$\\pm$   0.03&    1.75$\\pm$   0.02&   -1.62$\\pm$   0.13& 7238$\\pm$13&    0.64$\\pm$   0.22&    2.89$\\pm$   0.34&    4.80$\\pm$   0.04&    4.96$\\pm$   0.02\\\\ 010231.17-722134.0&   0.69$\\pm$   0.04&    1.83$\\pm$   0.03&   -2.13$\\pm$   0.15& 7107$\\pm$18&    0.91$\\pm$   0.18&    2.93$\\pm$   0.19&    5.46$\\pm$   0.06&    5.16$\\pm$   0.02\\\\ 005556.74-732133.6&   0.77$\\pm$   0.05&    1.77$\\pm$   0.03&   -1.80$\\pm$   0.17& 7193$\\pm$18&    0.70$\\pm$   0.21&    2.89$\\pm$   0.31&    5.00$\\pm$   0.06&    5.07$\\pm$   0.03\\\\ 003934.35-730433.9&   0.76$\\pm$   0.03&    1.79$\\pm$   0.02&   -1.92$\\pm$   0.13& 7161$\\pm$13&    0.76$\\pm$   0.21&    2.90$\\pm$   0.28&    5.16$\\pm$   0.04&    5.13$\\pm$   0.02\\\\ 010647.29-723053.7&   0.73$\\pm$   0.04&    1.81$\\pm$   0.02&   -2.03$\\pm$   0.13& 7124$\\pm$14&    0.77$\\pm$   0.21&    2.88$\\pm$   0.27&    5.31$\\pm$   0.04&    5.27$\\pm$   0.02\\\\ 005155.37-724909.5&   0.73$\\pm$   0.04&    1.77$\\pm$   0.02&   -1.78$\\pm$   0.17& 7183$\\pm$16&    0.65$\\pm$   0.23&    2.86$\\pm$   0.35&    5.00$\\pm$   0.05&    5.19$\\pm$   0.03\\\\ 010316.37-724816.2&   0.66$\\pm$   0.12&    1.90$\\pm$   0.03&   -2.56$\\pm$   0.13& 6968$\\pm$18&    1.14$\\pm$   0.12&    2.92$\\pm$   0.11&    6.20$\\pm$   0.15&    5.54$\\pm$   0.02\\\\ 010316.37-724816.2&   0.66$\\pm$   0.12&    1.90$\\pm$   0.03&   -2.56$\\pm$   0.13& 6968$\\pm$18&    1.14$\\pm$   0.12&    2.92$\\pm$   0.11&    6.20$\\pm$   0.15&    5.54$\\pm$   0.02\\\\ 005527.97-724136.8&   0.76$\\pm$   0.04&    1.81$\\pm$   0.03&   -2.03$\\pm$   0.16& 7118$\\pm$17&    0.75$\\pm$   0.22&    2.87$\\pm$   0.29&    5.32$\\pm$   0.05&    5.35$\\pm$   0.03\\\\ \\hline \\hline \\end{tabular} } \\end{table*}"
    },
    "0910/0910.3751_arXiv.txt": {
        "abstract": "Various theoretical and observational results have been reported regarding the presence/absence of net electric currents in the sunspots. The limited spatial resolution of the earlier observations perhaps obscured the conclusions. We have analyzed 12 sunspots observed from Hinode (SOT/SP) to clarify the issue. The azimuthal and radial components of magnetic fields and currents have been derived. The azimuthal component of the magnetic field of sunspots is found to vary in sign with azimuth. The radial component of the field also varies in magnitude with azimuth. While the latter pattern is a confirmation of the interlocking combed structure of penumbral filaments, the former pattern shows that the penumbra is made up of a ``curly interlocking combed\" magnetic field. The azimuthally averaged azimuthal component is seen to decline much faster than 1/$\\varpi$ in the penumbra, after an initial increase in the umbra, for all the spots studied. This confirms the confinement of magnetic fields and absence of a net current for sunspots as postulated by \\cite{parker96}. The existence of a global twist for a sunspot even in the absence of a net current is consistent with a fibril-bundle structure of the sunspot magnetic fields. ",
        "introduction": "Sunspots have shown evidence for twist even from the time of \\cite{hale25,hale27} who postulated the hemispheric rule for the chirality of chromospheric whirls. This was later confirmed with a larger data set by \\cite{rich41}. Evidence for photospheric chirality could be seen in early continuum images of sunspots, obtained with exceptional image quality. Later, photospheric vector magnetograms showed global twist inferred from the non-vanishing averages of the force-free parameter (\\cite{pcm94,hagi04,nand06} and references therein). The non-force-free nature of photospheric magnetic field in the sunspots, prompted \\cite{tiw09a} to propose the signed shear angle (SSA) as a more robust measure of the global twist of the sunspot magnetic field.  Although, the sign of SSA matches well with the sign of the global alpha parameter, the magnitudes are not so well correlated. The physical significance of a globally averaged $\\alpha$ parameter rests heavily on the existence of a net current in the photospheric sunspot magnetic field. One way of arriving at a global $\\alpha$ is by taking the ratio of total vertical current to the total flux (integral method). This value was found to agree with the values obtained by other methods \\citep{hagi04}.  For a monolithic sunspot magnetic field, the global twist and net current is expected to be well correlated by Ampere's Law. However, the existence of a net current is ruled out theoretically for fibril bundles as well as for monolithic fields with azimuthal field decreasing faster than 1/$\\varpi$, where $\\varpi$ is the radial distance from the spot center \\citep{parker96}. Several attempts to resolve this problem using vector magnetograms have not been very conclusive so far \\citep{wilk92,leka96,wheat2000}.  A resolution of this problem can be used to disentangle the relation between global twist and the global $\\alpha$ parameter. Also, the resolution is needed to evaluate the so called hemispheric helicity rule seen in the global $\\alpha$ parameter calculated from photospheric vector magnetograms \\citep{pcm94,pcm95,hagi04,nand06}. The availability of high resolution vector magnetograms from Hinode (SOT/SP), gives us the best opportunity so far to address this problem. The effect of polarimetric noise is expected to be negligible in the estimation of magnetic parameters \\citep{tiw09b} from these data.  In this Letter we obtain an expression for the net current using a generalization of the expression obtained by \\cite{parker96}. We then proceed to measure this current from several vector magnetograms of nearly circular sunspots. We finally discuss the results and present our conclusions. ",
        "conclusions": "It is well known for astrophysical plasmas, that the plasma distorts the magnetic field and the curl of this distorted field produces a current by Ampere's law \\citep{parker79}. Parker's (1996) expectation of net zero current in a sunspot was chiefly motivated by the concept of a fibril structure for the sunspot field. However, he also did not rule out the possibility of vanishing net current for a monolithic field where the azimuthal component of the vector field in a cylindrical geometry declines faster than 1/$\\varpi$. While it is difficult to detect fibrils using the Zeeman effect notwithstanding the superior resolution of SOT on {\\it Hinode}, the stability and accuracy of the measurements have allowed us to detect the faster than 1/$\\varpi$ decline of the azimuthal component of the magnetic field, which in turn can be construed as evidence for the confinement of the sunspot field by the external plasma. The resulting pattern of curl {\\bf B} appears as a drop in net current at the sunspot boundary.  If this lack of net current turns out to be a general feature of sunspot magnetic fields in the photosphere, then measurement of helicity from a global average of the force-free parameter becomes suspect. On the other hand, sunspots are evidently twisted at photospheric levels, as seen from the non-vanishing average twist angle as well as the SSA (Table 1). Although the existence of a global twist in the absence of a net current is possible for a monolithic sunspot field \\citep{baty00,arch04,fan04,aula05}, a fibril model of the sunspot field can accommodate a global twist even without a net current \\citep{parker96}.  The spatial pattern of current density in a sunspot (e.g., left panel of Figure 4) is really a manifestation of the deformation of the magnetic field ($\\nabla\\times\\bf B$) by the forces applied by the plasma. The Lorentz force exerted by the field on the plasma produces an equal and opposite force by the plasma, thereby confining the field. Thus our analysis actually shows the pattern of the forces exerted by the plasma on the field. The sharp decline of the azimuthal field with radial distance thus shows the confinement of the sunspot magnetic field by the radial gradient of the plasma pressure.  Theoretical understanding of the penumbral fine structure has improved considerably in recent times \\citep{thom02,weis04}. The onset of a convective instability for magnetic field inclination exceeding a critical value was proposed by \\cite{tilde03} and \\cite{hurl02}. A bifurcation in the onset \\citep{ruck95} could explain other features like hysteresis in the appearance of penumbra as a function of sunspot size. Numerical simulation of magneto-convection also steadily improved \\citep{hein07,remp09a}, culminating in very realistic production of penumbral field structure \\citep{remp09b}. It is possible, owing to the random and stochastic nature of convective structures, that no net twist in the simulated spot field would be produced by convection for negligible Coriolis force. If so, it would be very interesting to simulate magneto-convection in a twisted sunspot field. In this case, would the resulting fine structure mimic the observed ``curly interlocking combed'' structure of the penumbral magnetic field? If not, we must look elsewhere for explaining the ``curly interlocking combed'' structure. A twisted fibril bundle would then be a solution. Recent examples of filamentary penumbral structures based on such cluster models \\citep{sola93,spru06,scha06} have also been proposed.  \\cite{parker96} also mentions the possibility of net currents in the corona, continuing down to the height where the first cleaving takes place. It would therefore be imperative to look for net currents at higher reaches of the solar atmosphere. This is very important because several theories of flares \\citep{melr95} and CME triggers \\citep{forb91,klie06} rely heavily on the existence of net currents in the corona above the sunspots.  Future large ground based telescopes equipped with adaptive optics and multi spectral line capabilities would go a long way in addressing these issues. In the meantime, direct measurement of the global twist of sunspots using parameters like the SSA should serve as proxies for estimating the net currents of active regions in the corona. The SSA will also be a better parameter to base a fresh look at the hemispheric rule in photospheric chirality."
    },
    "0910/0910.1081_arXiv.txt": {
        "abstract": "We report the first detection of the Zeeman effect in the 36 GHz Class I \\methanol\\ maser line. The observations were carried out with 13 antennas of the EVLA toward the high mass star forming region M8E. Based on our adopted Zeeman splitting factor of $z = 1.7$ Hz \\mG, we detect a line of sight magnetic field of $-31.3 \\pm 3.5$ mG and $20.2 \\pm 3.5$ mG to the northwest and southeast of the maser line peak respectively. This change in sign over a 1300 AU size scale may indicate that the masers are tracing two regions with different fields, or that the same field curves across the regions where the masers are being excited. The detected fields are not significantly different from the magnetic fields detected in the 6.7 GHz Class II \\methanol\\ maser line, indicating that \\methanol\\ masers may trace the large scale magnetic field, or that the magnetic field remains unchanged during the early evolution of star forming regions. Given what is known about the densities at which 36 GHz \\methanol\\ masers are excited, we find that the magnetic field is dynamically significant in the star forming region. ",
        "introduction": "\\label{sINTRO}  Magnetic fields likely play an important role in the star formation process (e.g., \\citealt{tt2008}, and references therein). The nature of this role is not yet clear, however, primarily due to the scarcity of observational data on magnetic fields in star forming regions. The Zeeman effect remains the most direct method for measuring magnetic field strengths (e.g., \\citealt{rmc99}). Observations of the Zeeman effect in H I and OH thermal lines have revealed the strength of the magnetic field in the lower density envelopes of molecular clouds (e.g., \\citealt{bt01}). Observations of the Zeeman effect in \\water\\ masers, on the other hand, offer a window into magnetic fields in the highest density regions ($n \\sim 10^9~\\cm{-3}$) of star forming regions (\\citealt{sarma2008}; \\citealt{vdv06}).  Interstellar masers, being compact and intense, are extremely effective probes of young star forming regions. Recently, \\citet{wv2008} carried out the first systematic study of the Zeeman effect in the 6.7 GHz methanol maser line. Such (6.7 GHz) masers are examples of Class II methanol masers, which are known to be probes of the early phases of high mass star forming regions. Yet another class of masers, namely Class I methanol masers, may probe even earlier phases of star forming regions (\\citealt{pp2008}). This appears to be the case for M8E, which is a high mass star forming region located at $l = 6\\arcdeg.05$, $b = -1\\arcdeg.45$, at a distance of 1.5 kpc (\\citealt{scf1984}; \\citealt{gg76}). M8E is comprised of a bright infrared source (M8E-IR), and a radio \\ion{H}{2}\\ region about 8$\\arcsec$ to its northwest (\\citealt{scf1984}; \\citealt{dw2009}). Low spatial resolution CO observations by \\citet{wright77} showed that M8E is embedded in a dense molecular core. At higher resolution, \\citet{mhs92} mapped a bipolar outflow in the $^{12}$CO emission line (\\vLSR = 11 \\kmS), with blue and red wings spread over 50 \\kmS. Based on their VLTI observations, \\citet{linz2009} concluded that M8E-IR contains a central protostar of mass 10-15 \\Msun.  In this paper, we report the \\textsl{first} detection of the Zeeman effect in the 36 GHz Class I \\methanol\\ maser line. The observations were carried out with 13 antennas of the Expanded Very Large Array (EVLA) toward the high mass star-forming region M8E, and represent an early success of the EVLA. In \\S\\ \\ref{sODR}, we present details of the observations and reduction of the data. The analysis involved in extracting magnetic field information from the Zeeman effect is given in \\S\\ \\ref{sANAL}. The results are presented and discussed in \\S\\ \\ref{sR}. ",
        "conclusions": "\\label{sCONC} We have detected for the first time the Zeeman effect in the 36 GHz Class I \\methanol\\ maser line toward the high mass star forming region M8E. The detected values for the line-of-sight magnetic field are $-31.3 \\pm 3.5$ mG and $20.2 \\pm 3.5$ mG. There may be systematic bias in these values due to the assumed Zeeman splitting factor of $z = 1.7$ Hz \\mG, which is based on laboratory experiments on the 25 GHz \\methanol\\ maser line. Our detected values are similar to the magnetic fields detected in 6.7 GHz Class II \\methanol\\ masers by \\citet{wv2008}. This suggests that \\methanol\\ masers may trace the large-scale magnetic field; alternatively, the magnetic field may remain the same during the early evolution of the star forming region between the time when Class I and Class II masers are excited. Our detected values for \\Blos\\ imply that the magnetic field is dynamically significant to these regions.  Our detection gives rise to a host of questions that need further observational and theoretical work. We are planning future observations of the Zeeman effect in more 36 GHz masers and at higher resolution. We are also planning for observations of 44 GHz masers with the aim of detecting the Zeeman effect. As mentioned above, the 44 GHz line constitutes the other prominent Class I \\methanol\\ maser line. We are also planning to carry out high angular resolution observations of the Zeeman effect in regions containing both Class I and Class II masers, in order to compare the strength of magnetic fields between the regions traced by these classes. From theorists, we seek detailed models for Class I masers in order to understand their properties better, especially the range of densities at which they are excited. Also, there is a pressing need for laboratory studies of the Zeeman splitting coefficient for \\methanol\\ masers.  The resolution of some or all of these issues will open up yet another important window into the early stages of high mass star forming regions."
    },
    "0910/0910.4989_arXiv.txt": {
        "abstract": "We present axisymmetric hydrodynamical simulations of the long-term accretion of a rotating gamma-ray burst progenitor star, a ``collapsar,'' onto the central compact object, which we take to be a black hole.  The simulations were carried out with the adaptive mesh refinement code FLASH in two spatial dimensions and with an explicit shear viscosity.  The evolution of the central accretion rate exhibits phases reminiscent of the long GRB $\\gamma$-ray and X-ray light curve, which lends support to the proposal by \\citet{Kumar:08a,Kumar:08b} that the luminosity is modulated by the central accretion rate.  In the first ``prompt'' phase characterized by an approximately constant accretion rate, the black hole acquires most of its final mass through supersonic quasiradial accretion occurring at a steady rate of $\\sim 0.2~M_\\odot ~\\textrm{s}^{-1}$.  After a few tens of seconds, an accretion shock sweeps outward through the star.  The formation and outward expansion of the accretion shock is accompanied with a sudden and rapid power-law decline in the central accretion rate $\\dot M\\propto t^{-2.8}$, which resembles the $L_{\\rm X}\\propto t^{-3}$ decline observed in the X-ray light curves.  The collapsed, shock-heated stellar envelope settles into a thick, low-mass equatorial disk embedded within a massive, pressure-supported atmosphere, similar to the picture proposed by \\citet{Begelman:08} for ``quasistars.''  After a few hundred seconds, the inflow of low-angular-momentum material in the axial funnel reverses into an outflow from the surface of the thick disk. Meanwhile, the rapid decline of the accretion rate slows down, or even settles a in steady state with $\\dot M\\sim 5\\times10^{-5}~M_\\odot~\\textrm{s}^{-1}$, which resembles the ``plateau'' phase in the X-ray light curve.  While the duration of the ``prompt'' phase depends on the resolution in our simulations, we provide an analytical model taking into account neutrino losses that estimates the duration to be $\\sim 20~\\textrm{s}$.  The model suggests that the steep decline in GRB X-ray light curves is triggered by the circularization of the infalling stellar envelope at radii where the virial temperature is below $10^{10} ~\\textrm{K}$, such that neutrino cooling shuts off and an outward expansion of the accretion shock becomes imminent; GRBs with longer prompt $\\gamma$-ray emission have more slowly rotating envelopes. ",
        "introduction": "\\label{sec:intro} \\setcounter{footnote}{0}  Observations of long gamma-ray bursts (GRBs) carried out with the NASA {\\it Swift} satellite have shown that the $\\gamma$-ray prompt emission ceases after about a minute in the observer frame, corresponding to tens of seconds in the rest frame of the progenitor star.  The $\\gamma$-ray light curve, converted to a fiducial X-ray spectral band, smoothly joins the X-ray light curve, which declines rapidly, ($t^{-3}$ or faster) lasting for about $80$ to $300~\\textrm{s}$ \\citep{Tagliaferri:05,Nousek:06,OBrien:06}.  The rapid decline is often followed by a phase, from about $\\sim 10^3$ to $10^4~\\textrm{s}$, during which the X-ray flux is roughly constant or declines more slowly with time.  The X-ray light curves of some GRBs exhibit ``flares'' where the flux increases suddenly by a factor of $\\lesssim 10^2$ and drops precipitously, with the rise and decline associated with the flare occurring on a time scale much shorter than the age of the burst \\citep[see, e.g.,][]{Burrows:05,Falcone:06}.  Following about $\\sim10^3-10^4~\\textrm{s}$, a more rapid decline of the luminosity resumes \\citep[see, e.g.,][and references therein]{ZhangB:06}, and occasionally steepens further at $\\sim10^4-10^5~\\textrm{s}$ \\citep[e.g.,][]{Vaughan:06}.  The goal of the present work is to utilize two-dimensional hydrodynamic simulations to test the hypothesis \\citep{Kumar:08a,Kumar:08b} that this characteristic structure of the X-ray light curve, which was summarized by \\citet{ZhangB:06}, reflects a modulation in the rate of central accretion of a rotating progenitor star onto a black hole or a neutron star, as in the collapsar model of GRBs \\citep{Woosley:93,MacFadyen:99,MacFadyen:01,Woosley:06b}.  We do not attempt to explore the implications of the potential presence of a magnetosphere, as in the magnetar model for GRBs \\citep[e.g.,][]{Duncan:92,Wheeler:00,ZhangB:01,Thompson:04,Komissarov:07,Bucciantini:07,Bucciantini:09}. We will attempt to gain insight in the origin of the steady $\\gamma$-ray luminosity (the prompt phase which we will refer to as ``Phase 0''), the rapid decline in the X-ray light curve (Phase I in the nomenclature of \\citealt{ZhangB:06}), and phase of quasi-steady luminosity or slow decline (Phase II).  We will briefly attempt to extrapolate the results of our simulations to the subsequent steeper decline phases (Phases III and IV).  \\citet{Kumar:08a,Kumar:08b} obtained the key features of the $\\gamma$-ray and X-ray light curve by estimating central accretion rate resulting from the free (i.e., ballistic) infall of a rotating progenitor star.  In this picture, the material that has sufficient initial angular momentum to circularize outside of the innermost stable circular orbit (ISCO) of the black hole, forms a disk in the equatorial plane, and subsequently accretes via disk accretion \\citep{Narayan:01}.  If the luminosity is then assumed to be proportional to the central accretion rate, and if the distance of the $\\gamma$-ray or X-ray emitting region from the center of the star is assumed to be approximately independent of time on time scales $10-10^5~\\textrm{s}$, an accretion model directly translates into a synthetic light curve that can be compared with an observed light curve.  \\citet{Kumar:08a,Kumar:08b} have shown that with the simplest accretion model involving ballistic infall onto the midplane (assumption also made by \\citealt{Janiuk:08} and \\citealt{Cannizzo:09}) and subsequent disk accretion, the mapping of the mass accretion history onto the light curve provides a powerful insight into the stratification and angular momentum structure of the progenitor star.  In their ballistic infall model, \\citet{Kumar:08a,Kumar:08b} were not able to discriminate between models in which the quasi-steady activity in Phase II arose from disk accretion, or from late-time accretion from an extended stellar envelope. Departures from ballistic infall are expected if the infalling material passes through an accretion shock \\citep[see, e.g.,][]{MacFadyen:99,Lee:06,Nagataki:07,LopezCamara:09}, or if the disk launches a thermal \\citep[][]{MacFadyen:03a,MacFadyen:03b,Kohri:05} or magnetohydrodynamic \\citep[e.g.,][]{Proga:03c} outflow (``wind'') that can interfere with the infall.  The existence of the outflow is particularly interesting because of the potential for nucleosynthesis in the free neutron-rich outflow launched from the inner part of the disk \\citep[see, e.g.,][]{Pruet:03,Surman:06,Fujimoto:07,Nagataki:07,Maeda:09} and because of the potential that the outflow can deplete the accreting stellar envelope and limit the envelope mass that is accreted onto the central black hole.  During the first $\\sim10^2~\\textrm{s}$ following the formation of the central black hole when the accretion rate is $\\dot M\\gg 10^{-3}M_\\odot~\\textrm{s}^{-1}$ (the precise condition depends on the black hole spin and shear stress-to-pressure ratio $\\alpha$ in the disk), the inner accretion disk cools by neutrino emission and nuclear photodisintegration and accretes in a radiatively efficient fashion, except for in the very inner, optically thick region \\citep[e.g.,][]{Popham:99,Narayan:01,DiMatteo:02,Chen:07}.  Instabilities in the thin disk have been cited as a candidate class of mechanisms that could produce the observed X-ray flares \\citep{Perna:06,Lazzati:07,Lazzati:08} and could also produce detectable gravitational radiation \\citep{Piro:07}.  Our global axisymmetric models are the necessary stepping stone toward the substantially more computationally demanding three-dimensional simulations that will be required to pin down any nonaxisymmetric instabilities in the accreting collapsar \\citep[see, e.g.,][]{Rockefeller:06}.  We employ two-dimensional unmagnetized hydrodynamic simulations of the collapse, circularization, and accretion of a stellar envelope onto a central point mass, which we assume to be a black hole; relativistic corrections to the gravitational potential are ignored in our simulations since the innermost grid point lies at over $20$ Schwarzschild radii in the simulation extending to the smallest radius from the black hole.  The torque and dissipation arising from the $R-\\phi$ component of the magnetic stress is emulated with a Navier-Stokes term parameterized by an $\\alpha$-viscosity prescription. For comparison with the X-ray light curve, we measure the central accretion rate.  We track the flow of mass and energy at spherical radii $10^8~\\textrm{cm}\\lesssim r\\lesssim  10^{11}~\\textrm{cm}$ and interpret the results in view of the existing knowledge on radiatively-inefficient accretion flows.  We observe an outflow and measure the rate at which the accreting stellar envelope is lost to the outflow.   The mechanics of post-core-collapse accretion and outflows is key to estimating the final mass of the black hole and the nucleosynthetic composition of the ejected matter \\citep[e.g.,][and references therein]{Zhang:08}.  The method that we develop here can in future be utilized to estimate the masses of the black holes resulting from the collapse of massive, initially metal-poor ``Population III'' stars as well as from the collapse of the even more massive, hypothetical ``supermassive stars,'' in the presence of rotation.  In this work we do not simulate the neutrino-cooled disk, and in the simulations simply impose that the mass that crosses the innermost cylindrical radius of our simulation, $R_{\\rm min} = (0.5-2) \\times 10^{8} ~\\textrm{cm}$, is instantaneously incorporated inside the black hole and does not provide any further energetic feedback while at radii $R<R_{\\rm min}$.  This very rough assumption is bound to fail in general; it is most compatible with the regime in which the transition from efficient to inefficient cooling occurs at $R\\gtrsim R_{\\rm min}$.  Since the transitional radius for efficient neutrino cooling recedes inward with the increasing stress-to-pressure ratio $\\alpha$ for a given accretion rate \\citep{Chen:07}, the assumption that cooling is efficient within $\\lesssim 10^8~\\textrm{ cm}$ is valid for $\\alpha\\lesssim 0.01$.   We also ignore nuclear photodisintegration when temperature rises above $\\sim (5-10)\\times10^9~\\textrm{K}$; in reality, the photodissociation allows for some heating via the capture by free nucleons of the neutrinos emitted in the inner disk \\citep{Nagataki:07}, which we do not model.  However, we do incorporate neutrino cooling in a simple analytical model for the evolution of the accretion shock at the radii that we do not resolve, $\\lesssim 5\\times10^7~\\textrm{cm}$.  In combination with the simulations, the model provides a theory for the duration of the prompt emission phase observed in the $\\gamma$-rays.  \\citet{Cannizzo:09} speculate that a cool, thin disk may form at large radii ($R\\sim 10^{11}~\\textrm{cm}$) at the onset of Phase II, and attribute the structure of the X-ray light curve to the long-term evolution and slow central accretion of this extended disk.  We will see that the formation of the extended thin disk cannot be taken for granted due to the presence of a massive pressure-supported convective atmosphere around the inner disk.  This work is organized as follows.  In Section \\ref{sec:algorithm}, we discuss our numerical algorithm.  In Section \\ref{sec:results}, we present the results our simulations.  In Section \\ref{sec:discussion}, we present an analytical model for the neutrino-cooled central accretion that we do not resolve in the simulations, and provide a theory for the duration of the prompt accretion phase and the triggering of the steep decline of the X-ray light curve. We also attempt to extrapolate the evolution of the accretion rate beyond the duration of the simulations. Finally, in Section \\ref{sec:conclusions}, we summarize our conclusions. ",
        "conclusions": "\\label{sec:conclusions}  We have conducted hydrodynamic simulations of the viscous post-core-collapse accretion of a rapidly rotating $\\sim 14 ~M_\\odot$ Wolf-Rayet star of \\citet{Woosley:06a} onto the central black hole.  The axially-symmetric simulations were carried out for up to $2,000~\\textrm{s}$ and resolved the radii down to $5\\times 10^7~\\textrm{cm}$ where the collapsing stellar material circularizes around the black hole.  The evolution of the central accretion rate in the simulations resembles the evolution of the observed GRB X-ray luminosity, which lends support to the hypothesis \\citep{Kumar:08a,Kumar:08b} that the X-ray luminosity is proportional to the rate with which stellar material accretes onto the black hole.  We have identified three phases in the evolution of the accretion rate in our simulations, which appear to correspond to Phases 0 (the prompt phase), Phase I, and Phase II in the nomenclature of \\citet{ZhangB:06}.  In the initial phase that in the simulations lasts $37-52~\\textrm{s}$, the accretion of low-angular-momentum material is quasiradial for $r>5\\times10^7~\\textrm{cm}$ and occurs at quasi-constant rate of $\\sim0.2~M_\\odot~\\textrm{s}^{-1}$.  The end of this phase is marked by the formation of an accretion shock at the smallest resolved radii.  The shock immediately propagates radially outwards through the supersonically infalling stellar envelope.  Simultaneously with the formation and the outer movement of the accretion shock, the accretion rate drops suddenly and precipitously.  We argue that the somewhat late onset of the accretion shock is an artifact of our not resolving the innermost two decades in radius outside the black hole's gravitational radius.  We supplement the simulations with an analytical model of the innermost accretion disk not resolved in the simulations, and suggest that the accretion shock forms early, within a fraction of the first second of the formation of the black hole, as several published simulations of the innermost neutrino-cooled region have shown, but only starts to propagate outward after $20~\\textrm{s}$, when the material that is reaching the equatorial plane has enough angular momentum to circularize at radii where the virial temperature is below $\\sim 10^{10}~\\textrm{K}$ and the cooling by neutrino emission is suppressed.  During the second phase characterized by a steep decline $\\propto t^{-2.7}$ of the accretion rate that lasts $\\sim 500~\\textrm{s}$, the accretion shock sweeps through the star, but a supersonic accretion of the shocked fluid in the axial funnel region proceeds unabated, at least in our simulations where the funnel has not been heated to high temperatures by the relativistic jet.  The thick disk containing rotationally-supported and pressure-supported fluid is convective; a high-entropy outflow from the inner, rotationally-supported region follows the accretion shock on its traversal through the star but remains bound within the star and appears to form a large-scale circulation pattern.  The steepness of the accretion rate decline seems to be the consequence of a rapid hydrodynamic readjustment of the shocked, convective, and circulating stellar envelope.  The steep decline of the accretion rate slows down or stalls after $\\sim 600~\\textrm{s}$, which appears to reflect the settling of a fraction of the stellar envelope in the state of near-hydrostatic equilibrium.  The inner, rotationally supported thick disk contains $\\sim 1\\%$ of the mass of the unaccreted envelope and extends to $\\sim 3\\times10^9~\\textrm{cm}$.  The thick disk is surrounded by a much more massive, \\emph{pressure-supported} atmosphere, which acts as a mass supply to the thick disk. At no point do we find evidence for the extended thin disk envisioned by \\citet{Cannizzo:09}.  The fluid above and below the thick disk is mostly unbound and the simulations thus exhibit a form of a ``disk wind.''  We speculate that depletion of the envelope through accretion onto the black hole or mass loss in thermal outflows or winds could be responsible for the renewed steepening of the GRB X-ray light curve after $10^3-10^4~\\textrm{s}$.  More speculatively, the additional steepening of the light curve occasionally observed after $10^4-10^5~\\textrm{s}$ could be due to a pervasive thermal or radiatively-driven mass loss in the outer layers of the atmosphere."
    },
    "0910/0910.4171_arXiv.txt": {
        "abstract": "We present Chandra snapshot observations of the first large X-ray sample of optically identified fossil groups.  For 9 of 14 candidate groups, we are able to determine the X-ray luminosity and temperature, which span a range typical of large ellipticals to rich groups of galaxies.  We discuss these initial results in the context of group IGM and central galaxy ISM evolution, and we also describe plans for a deep X-ray follow-up program. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.4582_arXiv.txt": {
        "abstract": "The centers of ellipticals and bulges are formed dissipationally, via gas inflows over short timescales -- the ``starburst'' mode of star formation.  Recent work has shown that the surface brightness profiles, kinematics, and stellar populations of spheroids can be used to separate the dissipational component from the dissipationless ``envelope'' made up of stars formed over more extended histories in separate objects, and violently assembled in mergers.  Given high-resolution, detailed observations of these ``burst relic'' components of ellipticals (specifically their stellar mass surface density profiles), together with the simple assumptions that some form of the Kennicutt-Schmidt law holds and that the burst was indeed a dissipational, gas-rich event, we show that it is possible to invert the observed profiles and obtain the time and space-dependent star formation history of each burst.  We perform this exercise using a large sample of well-studied spheroids, which have also been used to calibrate estimates of the ``burst relic'' populations. We show that the implied bursts scale in magnitude, mass, and peak star formation rate with galaxy mass in a simple manner, and provide fits for these correlations.  The typical burst mass $M_{\\rm burst}$ is $\\sim10\\%$ of the total spheroid mass; the characteristic starburst timescale implied is a nearly galaxy-mass independent $t_{\\rm burst}\\sim10^{8}\\,$yr; the peak SFR of the burst is $\\sim M_{\\rm burst}/t_{\\rm burst}$; and bursts decay subsequently in power-law fashion as $\\dot{M}_{\\ast}\\propto t^{-2.4}$.  As a function of time, we obtain the spatial size of the starburst; burst sizes at peak activity scale with burst mass in a manner similar to the observed spheroid size-mass relation, but are smaller than the full galaxy size by a factor $\\sim10$; the size grows in time as the central, most dense regions are more quickly depleted by star formation as $R_{\\rm burst}\\propto t^{0.5}$.  Combined with observational measurements of the nuclear stellar population ages of these systems -- i.e.\\ the distribution of times when these bursts occurred -- it is possible to re-construct the dissipational burst contribution to the distribution of SFRs and IR luminosity functions and luminosity density of the Universe.  We do so, and show that these burst luminosity functions agree well with the observed IR LFs at the brightest luminosities, at redshifts $z\\sim0-2$. At low luminosities, however, bursts are always unimportant; the transition luminosity between these regimes increases with redshift from the ULIRG threshold at $z\\sim0$ to Hyper-LIRG thresholds at $z\\sim2$.  At the highest redshifts $z\\gtrsim2$, we can set strict upper limits on starburst magnitudes, based on the maximum stellar mass remaining at high densities at $z=0$, and find tension between these and estimated number counts of sub-millimeter galaxies, implying that some change in bolometric corrections, the number counts themselves, or the stellar IMF may be necessary.  At all redshifts, bursts are a small fraction of the total SFR or luminosity density, approximately $\\sim5-10\\%$, in good agreement with estimates of the contribution of merger-induced star formation. ",
        "introduction": "\\label{sec:intro}    Understanding the global star-formation history of the Universe remains an important unresolved goal in cosmology.  Of particular interest is the role played by mergers in driving star formation and/or the infrared luminosities of massive systems.  A wide range of observations support the view that violent, dissipational events (e.g.\\ gas-rich mergers) are important to galaxy evolution, and in particular that the central, dense portions of galaxy bulges and spheroids must be formed in such events; but less clear is their contribution to the global star formation process.  In the local Universe, the population of star-forming galaxies appears to transition from ``quiescent'' (undisturbed) disks -- which dominate the {\\em total} star formation rate/IR luminosity density -- at the luminous infrared galaxy (LIRG) threshold $10^{11}\\,\\lsun$ ($\\dot{M}_{\\ast}\\sim 10-20\\,\\msun\\,{\\rm yr^{-1}}$) to systems that are clearly merging and violently disturbed at a few times this luminosity. The most intense starbursts at $z=0$, ultraluminous infrared galaxies (ULIRGs; $L_{\\rm IR}>10^{12}\\,\\lsun$), are invariably associated with mergers \\citep[e.g.][]{joseph85,sanders96:ulirgs.mergers}, with dense gas in their centers providing material to feed black hole growth and to boost the concentration and central phase space density of merging spirals to match those of ellipticals \\citep{hernquist:phasespace,robertson:fp}.  Various studies have shown that the mass involved in these starburst events is critical to explain the relations between spirals, mergers, and ellipticals, and has a dramatic impact on the properties of merger remnants \\citep[e.g.,][]{LakeDressler86,Doyon94,ShierFischer98,James99, Genzel01,tacconi:ulirgs.sb.profiles,dasyra:mass.ratio.conditions,dasyra:pg.qso.dynamics, rj:profiles,rothberg.joseph:kinematics,hopkins:cusps.ell,hopkins:cores}.  At high redshifts, bright systems dominate more and more of the IR luminosity function \\citep[e.g.][]{lefloch:ir.lfs,perezgonzalez:ir.lfs,caputi:ir.lfs,magnelli:z1.ir.lfs}. Merger rates increase rapidly \\citep[by a factor $\\sim10$ from $z=0-2$; see e.g.][and references therein]{hopkins:merger.rates}, leading to speculation that the merger rate evolution may in fact drive the observed evolution in the cosmic SFR density, which also rises sharply in this interval \\citep[e.g.][and references therein]{hopkinsbeacom:sfh}.  However, many LIRGs at $z\\sim1$, and possibly ULIRGs at $z\\sim2$, appear to be ``normal'' galaxies, without dramatic morphological disturbances associated with the local starburst population or large apparent AGN contributions \\citep{yan:z2.sf.seds,sajina:pah.qso.vs.sf, dey:2008.dog.population,melbourne:2008.dog.morph.smooth, dasyra:highz.ulirg.imaging.not.major}.  At the same time, even more luminous systems appear, including large numbers of Hyper-LIRG (HyLIRG; $L_{\\rm IR}>10^{13}\\,\\lsun$) and bright sub-millimeter galaxies \\citep[e.g.][]{chapman:submm.lfs,younger:highz.smgs,casey:highz.ulirg.pops}. These systems exhibit many of the traits more commonly associated with merger-driven starbursts, including morphological disturbances, and may be linked to the emergence of massive, quenched (non star-forming), compact ellipticals at times as early as $z\\sim2-4$ \\citep{papovich:highz.sb.gal.timescales, younger:smg.sizes,tacconi:smg.maximal.sb.sizes, schinnerer:submm.merger.w.compact.mol.gas, chapman:submm.halo.clustering,tacconi:smg.mgr.lifetime.to.quiescent}.  In a series of papers, \\citet{hopkins:cusps.mergers, hopkins:cusps.ell,hopkins:cores} (hereafter H08c, H09b, H09e, respectively), the authors combined simulation libraries of galaxy mergers with observations of nearby ellipticals to develop a methodology to empirically separate spheroids into their two dominant physical components. First, a dissipationless component -- i.e.\\ an ``envelope,'' formed from the violent relaxation/scattering of stars already present in merging stellar disks that contribute to the remnant. Because disks are extended, with low phase-space density (and collisionless processes cannot raise this phase-space density), these stars will necessarily dominate the profile at large radii, hence the envelope, with a low central density.  Second, a dissipational or ``burst'' component -- i.e.\\ a dense stellar relic formed from gas which lost its angular momentum and fell into the nucleus of the remnant, turning into stars in a compact central starburst like those in e.g.\\ local ULIRGs (although these represent the most extreme cases).  Because gas radiates, it can collapse to very high densities, and stars formed in a starburst will dominate the profile within radii $\\sim0.5-1\\,$kpc, accounting for the high central densities of ellipticals.  The gas will reflect that brought in from merging disks, but could in principle also be augmented by additional cooling or stellar mass loss in the elliptical \\citep[see e.g.][]{ciottiostriker:recycling}.  In subsequent mergers, these two stellar components will act dissipationlessly, but the segregation between the two is sufficient that they remain distinct even after multiple, major ``dry'' re-mergers: i.e.\\ one can still, in principle, distinguish the dense central stellar component that is the remnant of the combined dissipational starburst(s) from the less dense outer envelope that is the remnant from low-density disk stars.  In H08c,  the methodology for empirically separating these components in observed systems is presented, and tested this on samples of nearby merger remnants from \\citet{rj:profiles}. Comparison with other, independent constraints such as stellar population synthesis models \\citep{titus:ssp.decomp, reichardt:ssp.decomp,michard:ssp.decomp} and galaxy abundance profiles \\citep{foster:metallicity.gradients.prep, mcdermid:sauron.profiles,sanchezblazquez:ssp.gradients}, as well as direct comparison of simulations with observed surface brightness profiles, galaxy shapes, and kinematics are used to demonstrate that the approach can reliably extract the dissipational component of the galaxy (see \\S~\\ref{sec:methods:obs}).  If this general scenario is correct, then the empirically identified burst relic components in local spheroids represent a novel and powerful new constraint on the history and nature of dissipational starbursts. In this paper, we present and develop these constraints -- and in particular, we show that they are not merely integral constraints on the masses and sizes of bursts. We show that the unique nature of dissipational star formation, together with the existence of some Kennicutt-Schmidt type relation between gas surface densities and star formation rates, means that the relic profiles can be {\\em inverted} to obtain the full time and radius-dependent star formation history of each galaxy starburst (\\S~\\ref{sec:methods}). We present the resulting, derived burst star formation histories for samples of hundreds of local, well-observed galaxy spheroids. We show that such starbursts follow broadly similar time-dependent behavior, with characteristic starburst star formation rates, durations, rise and decay rates, and spatial sizes that scale with galaxy mass and other properties according to simple scaling laws across $\\sim5$ decades in starburst and spheroid mass (\\S~\\ref{sec:results:properties}). Combining these inferred star formation histories (and resulting burst lightcurves) with empirical determinations of the ages of each starburst, we can reconstruct the starburst history of the Universe, including e.g.\\ the luminosity functions of such dissipational bursts at all redshifts, and their contribution to the global star formation rate density (\\S~\\ref{sec:results:history}). We discuss the uncertainties in this approach, compare with the previous, independent attempts to constrain these quantities via other observational methods, and summarize our results in \\S~\\ref{sec:discussion}.  Throughout, we adopt a $\\Omega_{\\rm M}=0.3$, $\\Omega_{\\Lambda}=0.7$, $h=0.7$ cosmology and a \\citet{chabrier:imf} stellar IMF, but these choices do not affect our conclusions. ",
        "conclusions": "\\label{sec:discussion}  It has long been believed that the centers of galaxy spheroids must be formed in dissipational starburst events, such that gas in the outer parts of a galaxy must lose angular momentum rapidly, and fall in on roughly a dynamical time to the galaxy center. Recently, observations have shown that it is possible to robustly separate the ``burst'' component of galaxy profiles from the outer, violently relaxed component owing to the scattering of progenitor galaxy stars formed over earlier, more extended periods.  As detailed, high-resolution observations of e.g.\\ spheroid stellar mass profiles, shapes, kinematics, and stellar population properties improve, such decompositions become increasingly robust and applicable to a wide range of systems.  In this paper, we use these observations as a novel, independent constraint on the nature of galactic starbursts. We stress that by ``starburst'' here, we do {\\em not} mean simply any high-SFR system; rather, we refer specifically to star formation in central, usually sub-kpc scale bursts owing to large central gas concentrations driven by angular momentum loss as described above.  These are commonly associated with galaxy-galaxy mergers, which have long been known to efficiently drive starbursts and violent relaxation \\citep[e.g.][]{lynden-bell67, toomre77,barnes.hernquist.91,barneshernquist96,barnes:review, mihos:starbursts.94,mihos:starbursts.96}. But they could also owe to any other sufficiently violent process; e.g.\\ gas inflows owing to disk bars or during dissipational collapse.  Regardless of origin, we can robustly identify the starburst relics and stellar mass surface density profiles in well-observed spheroids at $z=0$.  This methodology and tests of its accuracy have been discussed in \\citet{hopkins:cusps.mergers,hopkins:cusps.ell,hopkins:cores,hopkins:cusps.fp}. This alone is a powerful integral constraint on the star formation histories of these systems. But we show that they can in fact be used to provide more detailed information, by allowing for the inversion and recovery of the star formation history of each burst. We can thus use local observations to recover the full time-dependent and scale-dependent star formation history of each burst, if we couple these observed constraints to simple, well-motivated assumptions.  Namely, that some form of the Kennicutt-Schmidt relation (between star formation rate and gas surface density) applied in the formation of these stars, and that they formed in dissipational (i.e.\\ rapid and initially gas-rich, on these scales) events. Both assumptions are directly motivated by observations, but we also consider how uncertainties in them translate into uncertainties in the resulting constraints.  Performing this exercise, we recover a large number of empirically determined parameters of starbursts, and quantify how they scale as a function of mass and other properties.  This includes the total mass of the starburst and the time-dependent SFR. In general, we find that the constrained starbursts can be reasonably well-approximated by a simple power-law behavior (Equation~\\ref{eqn:mdot.time}), rising and decaying from a characteristic maximum SFR with characteristic half-life $t_{\\rm burst}$ and a typical late-time power-law slope of $\\dot{M}_{\\ast}\\propto (t/t_{\\rm burst})^{-\\beta}$ where $\\beta \\sim 2.0-3.0$.  The characteristic burst mass $M_{\\rm burst}$ is $\\sim10\\%$ of the total spheroid mass, but it scales weakly with galaxy mass in a manner similar to how disk galaxy gas fractions scale with their stellar masses (expected if disks are the pre-merger progenitors of spheroids), $M_{\\rm burst}\\approx 1/(1+[M_{\\ast}/10^{9.15}\\,\\msun]^{0.4})$.  Detailed discussion of this trend, and its consequences for the global structure and kinematics of spheroids, are presented in \\citet{hopkins:cusps.fp}. However, it already makes it clear that bursts should not dominate the SFR density.  They are a small fraction $\\sim10\\%$ of all stars in spheroids (let alone all stars).  That does not mean that bursts are unimportant, however.  It is clear that they control many of the properties of galaxies \\citep{cox:kinematics,robertson:fp, naab:gas,onorbe:diss.fp.details,ciotti:dry.vs.wet.mergers, jesseit:kinematics,jesseit:merger.rem.spin.vs.gas, covington:diss.size.expectation, hopkins:cusps.ell,hopkins:cores}, and they can account for the short-lived, highest-SFR systems in the Universe. Moreover, the scatter in burst mass is significant, $\\sim0.3-0.4\\,$dex, which is critical for explaining the most extreme starbursts in the Universe, which require both large absolute galaxy masses and gas fractions to reach the very high burst masses required.  The typical starburst timescale implied from the combination of observed surface densities and the \\citet{kennicutt98} law is a nearly galaxy-mass independent $t_{\\rm burst}\\sim10^{8}\\,$yr. Both the value of this timescale, and the weak scaling with galaxy and/or burst mass, agree well with the dynamical times in the central $\\sim$kpc of galaxies.  As above, though, there is considerable scatter of order $\\sim0.2\\,$dex.  Given such a short starburst timescale, and the relatively small total mass fractions involved in starbursts, it naturally follows that starbursts will represent only a small fraction of the star-forming galaxies at any stellar mass, at a particular instant. Given that the average number of bursts per galaxy is not large, the duty cycle of bursts should be $\\sim t_{\\rm burst}/t_{\\rm Hubble}$, or $\\sim1-5\\%$ from $z=0-3$.  Indeed, observations have shown that most galaxies at these redshifts lie on a normal star-forming sequence, without a large $\\sim1\\,\\sigma$ scatter from e.g.\\ merger-induced bursts \\citep{noeske:2007.sfh.part1, noeske:sfh,papovich:ssfr,bell:morphology.vs.sfr}. The small duty cycle here means that even if the burst SFR enhancement is large, it will not violate these constraints (appearing only at the $\\sim2-3\\,\\sigma$ level in the wings of the SFR distribution at a given mass).  These conclusions are supported by independent evidence from observational stellar population synthesis studies. Specific comparisons to the objects considered here, where available, are presented in detail in \\citet{hopkins:cusps.ell,hopkins:cores} and \\citep{foster:metallicity.gradients.prep}. It is well-established that constraints from abundances require the central portions of spheroids be formed in a similar, short timescale. And detailed decomposition of stellar populations into burst plus older stellar populations have yielded consistent results for the typical burst fractions and sizes \\citep{titus:ssp.decomp,schweizer:7252, reichardt:ssp.decomp,michard:ssp.decomp}.  As should be expected from the generic behavior above, bursts peak at SFRs of $\\sim M_{\\rm burst}/t_{\\rm burst}$, which follows $M_{\\rm burst}$ (and hence total spheroid mass) in a close-to-linear relation. The most massive local ellipticals -- especially those with total stellar masses of $\\gtrsim10^{12}\\,\\msun$ -- had extreme peak SFRs of $>1000\\,\\msun\\,{\\rm yr^{-1}}$.  Thus, it is at least possible that some local systems reached the highest SFRs inferred for massive, high-redshift starburst galaxies \\citep{papovich:highz.sb.gal.timescales, chapman:submm.lfs,walter:2009.hyperlirg.in.highz.qso.host}.  More moderate, but still massive ellipticals with $M_{\\ast}\\sim 1-5\\times10^{11}\\,\\msun$, reached a range of peak SFRs from $\\sim 30-500\\,\\msun\\,{\\rm yr^{-1}}$, corresponding to their forming fractions from $\\sim5-20\\%$ of their total masses in starbursts. These match well with the observed SFRs in more typical, local and $z\\sim1$ starbursts (in $\\sim L_{\\ast}$ galaxies), which have been specifically associated with mergers driving gas to galaxy centers, and forming the appropriate nuclear mass concentrations to explain e.g.\\ elliptical kinematics, sizes, phase space densities, and fundamental plane scalings \\citep{cox:kinematics, naab:gas,robertson:fp,jesseit:merger.rem.spin.vs.gas,hopkins:cusps.mergers, hopkins:cusps.fp,hopkins:cusps.evol}.  We similarly quantify starburst spatial sizes as a function of their mass, peak SFR, and time. Starburst sizes scale with starburst mass in a similar fashion as the spheroid mass-size relation, but are smaller than their host spheroids by a fraction similar to their mass fraction.  The most massive starbursts reach half-SFR (i.e.\\ half-light, in IR or mm wavelengths) size scales of $\\sim 1-5\\,$kpc. For typical spatial distributions, this implies total spatial extents of strong emission up to $\\sim10\\,$kpc.  These size scales also agree well with observations of the most massive high-redshift starbursting systems \\citep{younger:smg.sizes, tacconi:smg.maximal.sb.sizes, schinnerer:submm.merger.w.compact.mol.gas}, and of massive, compact ellipticals formed at high redshift, believed to be the relics of such starbursts \\citep[with, at that time, little envelope of dissipationless, low-density material yet accreted;][]{vandokkum:z2.sizes, cimatti:highz.compact.gal.smgs,trujillo:compact.most.massive, bezanson:massive.gal.cores.evol,hopkins:r.z.evol}.  For more typical, $\\sim L_{\\ast}$ starbursts, sizes range from $\\sim0.1-1\\,$kpc, also similar to those observed in merging systems \\citep{scoville86,sanders:review, hibbard.yun:excess.light,tacconi:ulirgs.sb.profiles,laine:toomre.sequence, rj:profiles}. The size difference between these and the most extreme objects follows from the much larger gas supply involved -- there is no dramatic difference in the relic starburst mass distribution {\\em shapes}.  Some claims have been made that the large sizes of high-redshift starbursts could imply that they are not scaled up analogues of local extreme starbursts; but we find here that they correspond naturally.  Scaling up a starbursting system in the starburst mass fraction will not preserve spatial size, but rather will scale along the starburst size-mass relation here, which appears to be smooth and continuous from the smallest starbursts with masses $\\sim10^{7}\\,\\msun$ to the largest with masses $>10^{11}\\,\\msun$.  The origin of the size-mass scaling is of considerable physical interest.  It has been proposed that in such starbursts the Eddington limit from radiation pressure on dusty gas sets a universal maximum central surface density, over all mass scales, from which this follows \\citep{hopkins:maximum.surface.densities}.  The important constraint, from our analysis, is that there is no discontinuity, and we provide the scaling relations that any such model must satisfy.  Combining these constraints with observational measurements of the nuclear stellar population ages of these systems -- i.e.\\ the distribution of times when these bursts occurred -- we show that it is possible to re-construct the dissipational burst contribution to the distribution of SFRs and IR luminosity functions and luminosity density of the Universe.  We show that the burst luminosity functions agree well with the observed IR LFs at the brightest luminosities, at redshifts $z\\sim0-2$. At low luminosities, however, bursts are always unimportant, as expected from their short duty cycles, noted above. Although the burst luminosity functions rise with redshift, they always represent low space densities, and the overall LF evolves rapidly. As such, the transition luminosity above which bursts dominate the IR LFs and SFR distributions increases with redshift from the ULIRG threshold at $z\\sim0$ to Hyper-LIRG thresholds at $z\\sim2$. This appears to agree well with recent estimates of the transition between normal star formation and mergers, along the observed luminosity functions.  Systematic morphological studies at low redshifts \\citep{sanders96:ulirgs.mergers} yield the conventional wisdom that -- locally -- the brightest LIRGs and essentially all ULIRGs are merging systems (see also references in \\S~\\ref{sec:intro}).  At high redshifts, similar studies have now been performed \\citep[see e.g.][and references therein]{tacconi:smg.mgr.lifetime.to.quiescent}.  They too find that the brightest sources are almost exclusively mergers, but with a transition point an order-of-magnitude higher in luminosity.  Other morphological studies at intermediate redshifts $z\\sim0.4-1.4$ have reached similar conclusions \\citep{bridge:merger.fractions}.  Integrating these luminosity functions, we find the burst contribution to the SFR and IR luminosity densities of the Universe, and show that it is small at all redshifts, rising from $\\sim1-5\\%$ at $z\\sim0$ to a roughly constant $\\sim4-10\\%$ at $z>1$.  This agrees well with recent attempts to estimate the contribution to the SFR density at $z=0-1$ specifically from merger-induced starbursts, using either pair or morphologically-selected samples \\citep{jogee:merger.density.08, robaina:2009.sf.fraction.from.mergers}.\\footnote{Note that it is important here to distinguish estimates of the SFR {\\em induced} by mergers, i.e.\\ that above what some control population would exhibit, from that simply in ongoing/identifiable mergers (since many criteria identify mergers for a timescale $\\sim$Gyr, much longer than the burst timescale). For example, a merger fraction of $10\\%$ would imply at least $10\\%$ of star formation in ongoing mergers, even if there were no starbursts and those systems were only forming stars at the same rate as they would in isolation.}  Given the completely independent nature of the constraints, and significant uncertainties involved in both, the agreement is good.  The small value is what is expected, given our previous determination that the typical burst mass is just $\\sim10\\%$ in $\\sim L_{\\ast}$ spheroids (the galaxies that dominate the stellar mass density).  But it clearly rules out merger-induced bursts driving the SFR density evolution of the Universe.  At the highest redshifts $z\\gtrsim2$, we can put strict upper limits on starburst intensities, based on the maximum stellar mass remaining at high densities at $z=0$, and find some tension between these and estimated number counts of sub-millimeter galaxies from \\citet{chapman:submm.lfs}.  This implies that some change may be necessary in either the number counts themselves, the bolometric corrections used to convert these observations to total IR luminosities, or the stellar IMF used to convert between SFR and IR luminosity.  However, the observations remain considerably uncertain, with the bolometric corrections relying sensitively on assumed dust temperatures, and the number counts subject to significant cosmic variance \\citep[see e.g.][]{austermann:2009.aztec.submm.source.counts}.  More observations, at new wavelengths and in larger, independent fields, are needed to resolve these discrepancies.  We compare our constraints on these histories with recent predictions from galaxy formation models and simulations, and find reasonably good agreement.  Both exhibit similar tension with estimates of the high-redshift, extremely luminous number densities. However, the models are able to match the inferred number densities presented here {\\em without} any change to the stellar IMF or requiring other exotic physics.  The systematic uncertainties in the models, especially at high redshift, are large, however, so the empirical constraints presented here provide a powerful new means of constraining the models and their input parameters.  For example, if models are constrained via e.g.\\ a halo occupation approach or otherwise, so as to match the observed merger fractions of galaxies as a function of redshift, then these numbers become relatively large ($>10\\%$) at redshifts $z\\gtrsim2$ \\citep[see e.g.][]{bundy:merger.fraction.new,conselice:mgr.pairs.updated, kartaltepe:pair.fractions,lin:mergers.by.type,bluck:highz.merger.fraction, hopkins:merger.rates,jogee:merger.density.08, bridge:merger.fraction.new.prep}. The most common assumption in many analytic and semi-analytic models is that, in such a major merger, the entire galaxy gas supply is channeled into the starburst. This, however, coupled with the high observed merger fractions (and implied merger rates), would lead to an over-prediction of the burst contribution to the SFR density, relative to our constraints here. For example, the predictions of such a model are given in \\citet{hopkins:merger.lfs}, and rise to $\\sim20-50\\%$ of the SFR density at $z>1-2$.  Recently, however, \\citet{hopkins:disk.survival} have pointed out that the processes that lead to angular momentum loss by the gas -- driving bursts in the first place -- rely on the stellar component of merging systems, and so become less efficient as gas fractions increase. Including these physics in the models (as the current generation to which we compare does) leads to an asymptotic maximum contribution to the SFR density similar to that found here, while still giving a good match to observed merger fractions.  Despite the simple nature of the assumptions involved in our analysis, we show in high-resolution hydrodynamic simulations that they work well in allowing us to recover the star formation histories of bursts (even where the simulated system is much more complex).  Especially in a statistical sense, our numerical experiments suggest that this approach is robust to a variety of detailed deviations from our idealized assumptions.  Testing this (even in a few objects) via high-resolution reconstruction of stellar populations in bursts, directly from their observed spatially-resolved SEDs, would provide a powerful test of this.  Insofar as we have compared with different simulations and models, the dominant uncertainty in our analysis is the nature -- in particular the slope, relevant for the extrapolation to high gas surface densities -- of the \\citet{kennicutt98} law.  We have considered a range of slopes suggested by different observations. Although they yield similar behavior at moderate and low SFRs, after approximately $\\sim10^{8}$\\,yr from the beginning of a burst, the predicted behavior at very early times in the bursts is quite different.  A larger \\citet{kennicutt98} relation index implies higher peak SFRs and more sharply peaked starbursts, increasing the predicted number counts of the most luminous sources and making it relatively more easy for models to account for the most extreme star-forming systems.  It therefore remains of considerable importance for observations to probe both gas density measurements and full SFR indicators in extreme systems, at low and high redshifts.  Finally, we note that our analysis is only possible because, as indicated by basic dynamics, simulations, and observations of stellar populations, starburst components of spheroids were formed in dissipational (i.e.\\ gas-rich, at their centers), rapid star forming events. The dissipationless ``envelopes'' surrounding the central, dense components in spheroids were not formed in such a manner (again indicated by both their structural and kinematic properties, simulations of their formation, and stellar population observations).  Rather, they represent the debris of stars from disks which were formed pre-merger, and assembled (and violently relaxed) dissipationlessly.  These stars were formed over extended periods of time, with new gas accretion onto the disk fueling new or continuous star formation \\citep[e.g.][]{keres:cooling.clumps.from.broken.filaments}, as opposed to a single massive inflow.  Especially in the most massive systems, they are also assembled from multiple systems, via e.g.\\ minor mergers contributing tidal material to the extended ``wings'' well-known in massive galaxies.  As such, there is not a straightforward means to invert their surface stellar mass density profiles to obtain their star formation history.  Such constraints will depend on other methods, such as e.g.\\ direct stellar population analysis. It should be borne in mind that this represents $\\sim90\\%$ of the mass in most spheroids, and so understanding star formation in disks remains critical to understanding the origin of stellar populations in spheroids."
    },
    "0910/0910.1472_arXiv.txt": {
        "abstract": "{}{We compile and analyze an extended database of light curve parameters scattered in the literature to search for correlations and study physical properties, including internal structure constraints.} {We analyze a vast light curve database by obtaining mean rotational properties of the entire sample, determining the spin frequency distribution and comparing those data with a simple model based on hydrostatic equilibrium.} {For the rotation periods, the mean value obtained is 6.95 h for the whole sample, 6.88 h for the Trans-neptunian objects (TNOs) alone and 6.75 h for the Centaurs. From Maxwellian fits to the rotational frequencies distribution the mean rotation rates are 7.35 h for the entire sample, 7.71 h for the TNOs alone and 8.95 h for the Centaurs. These results are obtained by taking into account the criteria of considering a single-peak light curve for objects with amplitudes lower than 0.15 mag and a double-peak light curve for objects with variability $>$0.15mag. We investigate the effect of using different values other than 0.15mag for the transition threshold from albedo-caused light curves to shape-caused light curves. The best Maxwellian fits were obtained with the threshold between 0.10 and 0.15 mag. The mean light-curve amplitude for the entire sample is 0.26 mag, 0.25 mag for TNOs only, and 0.26 mag for the Centaurs. The Period versus B-V color shows a correlation that suggests that objects with shorter rotation periods may have suffered more collisions than objects with larger ones. The amplitude versus H$_v{}$ correlation clearly indicates that the smaller (and collisionally evolved) objects are more elongated than the bigger ones.} {From the model results, it appears that hydrostatic equilibrium can explain the statistical results of almost the entire sample, which means hydrostatic equilibrium is probably reached by almost all TNOs in the H range [-1,7]. This implies that for plausible albedos of 0.04 to 0.20, objects with diameters from 300 km to even 100 km would likely be in equilibrium. Thus, the great majority of objects would qualify as being dwarf planets because they would meet the hydrostatic equilibrium condition. The best model density corresponds to 1100 kg/m$^3$.} ",
        "introduction": "The belt of remnant planetesimals in the so-called Kuiper belt, whose existence was confirmed observationally in 1992 \\citep{Jewitt1993}, is an important source of knowledge about the formation and early evolution of the Solar System. Features in the orbital distribution of the belt can provide important information about the early dynamical processes such as planetary migration, which was merely untested theories around 25 years ago \\citep{Fernandez1984} that now can be regarded as being well proven. The detailed architecture of the Kuiper belt might also be compatible with a very unstable phase when Jupiter and Saturn entered into resonance \\citep{Tsiganis2005, Morbidelli2005, Gomes2005} around 800Myr after the Solar System formation. Although this particular dynamical scenario is still disputed, there is no question that the architecture of the transneptunian belt can offer invaluable clues about the dynamical mechanisms acting in the early Solar System. In addition to dynamical inferences from the orbital characteristics of the TNOs, much can be learnt about the physical processes that took place in the solar nebula by means of the study of the physical properties of TNOs and related groups of objects. Since the TNOs are understood to be mixtures of ices and rocks from which the Centaurs and ultimately the Jupiter family comets (JFCs) originate \\citep{Jewitt2004b}, ground and space-based observations of all these populations can give us clues about the physics of the outer Solar System. From the ``Nice\" dynamical model \\citep{Tsiganis2005, Morbidelli2005, Gomes2005}, it is also likely that the Jupiter trojans are closely related to the TNOs \\citep{Morbidelli2005}, but this is not yet well established.  These objects are regarded as the least evolved objects in the Solar System, but they have experienced many different types of physical processes that have altered them considerably. These processes range from collisions, to space weathering. The collisional processes have probably left clear imprints in both the spin distribution and in the light-curve amplitude distribution of TNOs and Centaurs. The currently accepted theory based on the study of \\cite{DavisFarinella1997} on collisional evolution of the TNOs is that objects of sizes larger than 100 km are very slightly collisionally evolved, whereas the smaller fragments have experienced more collisional events and hence should have their rotations and shapes more significantly altered. This picture is also consistent with newer collisional evolution models \\citep{Benavidez2009}. However, detailed modelling of the evolution of the spin rates related to collisional models has not yet been carried out for the transneptunian belt.  Due to the dedicated efforts of two main groups studying short-term variability \\citep{Gutierrez2001, Ortiz2003a, Ortiz2003b, Ortiz2004, Ortiz2006, Ortiz2007, Moullet2008, Duffard2008,Sheppard2007, Sheppard-Jewitt2002, Sheppard2000} and occasional contributions by other authors, a database of about 40 rotation periods and 80 light-curve amplitudes could be extracted from the literature in 2008. Thus, we now have sufficient numbers of TNOs light curves to start analyzing them from a statistical point of view. In a review by \\cite{Sheppard2008}, those data were compiled and some conclusions were outlined. After that compilation, a new set of light curves from 29 TNOs and Centaurs became available \\citep{Thirouin2009} and those authors also presented evidence of possible biases that can be present in the light curve database gathered from the literature.  We also developed a simple model to interpret rotation rates and light-curve amplitudes together, to obtain information about densities and tensile strengths. Instead of analyzing each body separately and trying to retrieve information about its density and cohesion by using its spin rate and light-curve amplitude (and an assumed viewing angle), we address the mean internal properties by analyzing the population as a whole. In other words, we developed a model that can simultaneously interpret the distribution of spin rates and the distribution of light-curve amplitudes.  In this paper, light curve parameters of a vast sample are compiled and analyzed. We focus on the analysis of the period and shape distribution presented in Sect. 2. Then, with the current data, we complete a correlation analysis with several orbital parameters in Sect. 3. Internal properties of the TNOs are derived by applying a simple model in Sect. 4, a discussion of the results is presented in Sect. 5. Finally, the conclusions are presented in Sect. 6. ",
        "conclusions": "We have analyzed the light-curve data of \\cite{Thirouin2009} and others in the literature to derive light-curve parameters, which have been compiled here. We have presented the most complete distribution of rotation periods and amplitudes so far. For the rotation periods, the mean value obtained by a Maxwellian fit in rotational frequencies is 7.71 hours (3.11 cycles/day) considering only the TNOs, 8.95 hours (2.68 cycles/day) considering only Centaurs, and 7.35 hours (3.19 cycles/day) considering all the sample altogether. These results were obtained by taking into account the criteria of considering a single peak light curve for those objects with amplitudes lower than 0.15 mag and a double peak light curve for the more elongated ones.  This mean rotational period is slightly longer than the mean value for the largest Main Belt asteroids, 6.9 hours \\citep{Binzel1989}, but we state that this value may be affected by observational biases. It is important to stress that in photometry devoted to the study of the rotational properties of minor bodies, there exists and important observational bias towards shorter periods and higher amplitude light curves. How this bias can be estimated remains an open problem, but it certainly depends on the number of objects for which a period is obtained. Some steps in this direction were already taken by \\cite{Thirouin2009}. Concerning other biases, if observational time for larger telescopes is granted then smaller TNOs will be accessible for the determination of the light curves, and the sample will not only increase in size but the bias in size will become less important.  One of the most important correlations found in this work is the light curve amplitude and absolute magnitude of all the sample of KBOs. This correlation is not a surprise and is probably related to the most collisionally evolved small population, which should be more elongated objects. On the other hand, the less collisionally evolved population (larger and then brighter objects) is the more numerous one in our sample ($\\sim$60\\%), so we caution the reader that some of our results could be somewhat biased. Another interesting correlation is the rotational period and B-V color, which could indicate the age of the surface (younger surfaces being bluer). This correlation appears to be consistent with the idea that the more collisionally evolved objects spin faster and have bluer surfaces.  A simple shape model has been developed that is able to reproduce some results on the relative percentage of presence of MacLaurin/Jacobian figures, when assuming hydrostatic equilibrium. In the entire sample, less than 12\\% of the objects in hydrostatic equilibrium have a Jacobi shape. Most of the objects have MacLaurin figures that have small light curve amplitudes (because of albedo marks). The model that best fits the observations requires densities between 1000 to 1200 kg/m$^3$. We must mention that as was highlighted in \\cite{Thirouin2009} there is an observational bias in the percentage of low amplitude light curves, which means that perhaps 75\\% of the objects have low light curve amplitudes. This implies that the mean density is likely somewhat higher, around 1500 kg/m$^3$ (see Fig. \\ref{fig08}).  For plausible albedos of 0.04 to 0.20, the absolute magnitude range H=[-1,7] includes objects with diameters as small as $\\sim$120 km. Those objects would also be in hydrostatic equilibrium. We end with the provocative thought that the great majority of observed TNOs would qualify to be dwarf planets, because they appear to meet the hydrostatic equilibrium condition."
    },
    "0910/0910.0018_arXiv.txt": {
        "abstract": "The low-mass X-ray binary microquasar GRO~J1655--40 is observed to have a misalignment between the jets and the binary orbital plane. Since the current black hole spin axis is likely to be parallel to the jets, this implies a misalignment between the spin axis of the black hole and the binary orbital plane. It is likely the black holes formed with an asymetric supernova which caused the orbital axis to misalign with the spin of the stars. We ask whether the null hypothesis that the supernova explosion did not affect the spin axis of the black hole can be ruled out by what can be deduced about the properties of the explosion from the known system parameters. We find that this null hypothesis cannot be disproved but we find that the most likely requirements to form the system include a small natal black hole kick (of a few tens of $\\rm km\\, s^{-1}$) and a relatively wide pre-supernova binary. In such cases the observed close binary system could have formed by tidal circularisation without a common envelope phase. ",
        "introduction": "It has long been known that neutron stars have much greater space velocities than their likely progenitors \\citep{GO70} and it is now widely accepted that that this is because they acquire large velocity kicks in the supernova explosions in which they form \\citep{S70,S78}. The reasons for these kicks is still a matter for debate with the leading candidates being asymmetric neutrino emission and/or asymmetric mass release during the supernova explosion and core collapse \\citep{BP95,Pod02}.  We address here the question of the extent to which similar kicks may be present when black holes form. Because the material forming the black holes passes through the event horizon it is quite possible that few neutrinos can escape \\citep{gour93} so that the black hole could form with little or no natal kick. On the other hand, asymmetric collapse to form a black hole might lead to copious emission of gravitational waves \\citep{bonnell95, kobay03}. In addition, \\cite{LL94} note that in a binary system a black hole can form by accretion of matter on to a neutron star. In that case the kick given to the original neutron star would appear as a kick given to the current black hole. Evidence that black holes are indeed kicked comes from the work of \\cite{jonker04}. They looked at the out-of-plane distributions of low mass X-ray binaries and neutron stars. They found no significant difference between the two, leading to the conclusion that black holes are subject to similar kicks at their formation.   In this paper we accept the evidence that stellar black holes do indeed acquire a velocity kick when they form, but we inquire further into the nature of the kick. In particular we ask whether the mechanism which gives rise to the kick might also give rise to a misalignment of the black hole spin axis with the original spin axis of the star from which it formed. The black hole progenitor star most likely had its spin aligned with the binary orbit before the supernova. A simple spherical collapse would preserve the spin axis, but a more complicated collapse might not. We consider here the simple null hypothesis that the spin axis remains unaltered by the kick process and test the extent to which this might be contradicted by the evidence.  We focus on the microquasar GRO J1655-40. As we discuss in Section~\\ref{GRO}, there is considerable information for this system, about the size of the natal velocity kick and also on the misalignment between the current black hole spin and the binary orbital axis. In Section~\\ref{model} we outline the dynamics of natal kicks and their implications for spin/orbital misalignment and in Section~\\ref{apply} apply these results to GRO J1655-40. We discuss our results in Section~\\ref{conclude}. ",
        "conclusions": "\\label{conclude}  We have considered what can be learned about the natal kick acquired by a black hole in a supernova by considering the system GRO~1655--40. In line with the analysis of \\cite{Willems05} we make use of kinematic data, but we add the additional constraint that the current black hole spin is not aligned with the orbital rotation. Of course, if the supernova explosion and collapse process in which the black hole is formed gives not just a linear impulse, but also angular momentum, to the hole, then the current misalignment is just a measure of the added angular momentum. The task we set ourselves is to ask if it is possible, given the various constraints, to rule out the possibility that the natal kick (if any) added just linear, and not angular, momentum. To simplify the analysis we have fixed the masses of the stars before and after the supernova.  To be fully consistent we could allow these masses to vary too and assign probabilities based on the likelihood of a given set of parameters.  However one set of masses proved to be sufficient to not rule out the null hypothesis.  Our main results are summarised in Table~\\ref{xray}. We have considered randomly oriented natal kicks with Maxwellian velocity distributions of $265$ and $26.5\\,\\rm km \\,s^{-1}$ for various separations of the pre-supernova binary, parametrised in terms of its relative orbital velocity $v_{\\rm orb}$.  Column~2 gives the probability $P_1$ that a bound system with the appropriate properties can be formed, ignoring any information about misalignment. It is immediately evident, that for low-velocity kicks the probability of forming the system is high, being typically around a few tens of percent \\citep[compare][]{Willems05}. The probabilities are lower for higher-velocity kicks, being typically around a percent. That is only around one in a hundred such pre-supernova systems would end up looking like GRO~1655--40. This is still probably not unreasonable.  The remaining columns in Table~\\ref{xray} show the probabilities, given $P_1$ that the additional constraint of the misalignment angle $i$ can be satisfied, given our null hypothesis. It is evident that we are not able to rule out this null hypothesis, and so it is quite possible that the black hole formation process does not affect the angular as well as the linear momentum of the resulting black hole.  It is worth noting, however, that even without the alignment information, it is easier to form the observed system with a low-velocity kick.  Moreover, if the formation process did just impart linear momentum, then, with account for the fact that the likely alignment timescale is comparable to the age of the system since the formation of the hole \\citep{martin2008} so that the initial misalignment angle $i$ would have been in the range $20^\\circ - 40^\\circ$, then it is evident that there is a strong preference for the pre-supernova system to have be fairly wide. Note that $v_{\\rm orb}=50\\,\\rm km\\,s^{-1}$ corresponds to a binary separation of around $4.6\\,\\rm AU$ before the supernova. It is interesting that such a wide system could have avoided the traditional common envelope required to shrink the orbit \\citep{verbunt96}   We also note, that if the current misalignment angle were large (say, $i > 40^\\circ$) then it would become increasingly difficult to satisfy the constraints for GRO~1655--40. In this respect we note that V4641 Sgr is a microquasar similar in many respects to GRO~1655--40 and it has a current misalignment angle of $i \\approx 55^\\circ$ \\citep{O01}. If further observations were able to provide a constraint on its post-supernova systemic velocity, $v_{\\rm sys}$ then it might be able to comment more securely on the null hypotheses. For example if it could be shown to have a large space $v_{\\rm sys}$ then satisfying the null hypothesis might be problematic."
    },
    "0910/0910.5910_arXiv.txt": {
        "abstract": "{Neutral atomic hydrogen (\\hi) traces the interstellar medium (ISM) over a broad range of physical conditions. Its 21-cm emission line is a key probe of the structure and dynamics of galaxies. This line comprises very different temperature and density regimes on all scales, from tens of astronomical units to kiloparsecs. \\hi is the key element to study the evolution of the ISM in detail. To understand the physics of the ISM and to analyze the interplay between different phases it is mandatory to cover observationally a broad range of scales. But large scale imaging of galaxies, resolving at the same time all these scales is difficult; spatial resolution, as well as sensitivity and the field of view are currently rather limited.  The observational situation is much more favorable if we consider our own galaxy. Two major all sky 21-cm line surveys of the Milky Way will become available soon. The Galactic All Sky Survey (GASS) obtained with the Parkes 64-m telescope\\thanks{The Parkes Radio Telescope is part of the Australia Telescope which is funded by the Commonwealth of Australia for operation as a National Facility managed by CSIRO}~ for the southern hemisphere with a resolution of 16 arcmin is close to completion. The northern extension, the Effelsberg Bonn HI Survey (EBHIS)\\thanks{Based on observations with the 100-m telescope of the MPIfR (Max-Planck-Institut f\\\"ur Radioastronomie) at Effelsberg}~ with 9 arcmin resolution, will be available in 2010/2011; we refer to the talks by Kerp and Winkel. Here we discuss briefly the GASS and demonstrate the unprecedented quality of this survey.  The Galactic single dish 21-cm line surveys prepare the ground for future high resolution imaging of the Galactic \\hi distribution. Using the available short spacing informaton, the Australian Square Kilometre Array Pathfinder (ASKAP) will be capable to generate a truly panoramic view of the Milky Way HI gas distribution with arcsecond resolution for all declinations $ < 30\\deg$. Data from the Widefield ASKAP L-band Legacy All-Sky Blind Survey (WALLABY, see talk by Staveley-Smith) can be used to generate high resolution all sky maps. In comparison to the currently available interferometric International Galactic Plane Surveys (IGPS) the sensitivity will improve by a factor of 10. Most important is the all sky coverage which will overcome the rather limited spatial coverage of a few degrees around the Galactic plane for the IGPS.}  \\FullConference{Panoramic Radio Astronomy: Wide-field 1-2 GHz research on galaxy evolution\\\\ June 2-5 2009\\\\ Groningen, the Netherlands}  \\begin{document} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.5121_arXiv.txt": {
        "abstract": " ",
        "introduction": "Large redshift surveys are typically completed by observing with a Multi Object Spectrograph (MOS), obtaining spectra for many hundreds of sources simultaneously over large fields of view. The problem of how to optimise observing strategies to target sources distributed over some survey area with a given MOS, defining a field of view and number of simultaneous targets, falls into the ``area packing\" class of problems. Much work outside of astronomy has been devoted to such problems \\citep{megi84} which are usually intractable in a formal, provably-optimal, sense. In the case of the Anglo-Australian Telescope's (AAT) largest survey to date, the 2 degree Field Galaxy Redshift Survey \\citep[2dFGRS,][]{coll01}, the survey was created in a manner that minimised field overlaps in order to maximise area (the target magnitude limit being $b_{j}=19.45$). This obviously had an impact on the target completeness, and the observations  had to be weighted in order to account for the local levels of incompleteness. At the other extreme is the 6 degree Field Galaxy Survey \\citep[6dFGS,][]{jone04} that aimed for high levels of completeness within the local universe. In this case the filamentary structures (i.e.\\ non-uniform overdensities) present on small scales necessitate extremely non-uniform tile coverage and potentially large amounts of overlap among tiles, target densities varying from 6 to 30 galaxies per deg$^{2}$. Hence the optimal strategy for tiling is closely linked to the scientific objectives of the survey, and a generic approach will not be appropriate for all requirements.  Fibre fed MOS instruments typically have a circular field of view (FOV), as seen for example in the 2 degree Field \\citep[2dF,][]{lewi02}, 6 degree Field \\citep[6dF,][]{jone04}, Sloan Digital Sky Survey (SDSS) Spectrograph \\citep{york00}, Hectospec \\citep{fabr05} and Hydra \\citep{bard93}. Also typical is for survey regions to be rectangular in spherical coordinate geometry: recent examples include the 2dFGRS, Sloan Digital Sky Survey \\citep[SDSS,][]{abaz09} and Millennium Galaxy Catalogue \\citep[MGC,][]{lisk03,driv05}. This latter commonality is due to a number of allying factors: imaging CCDs used for input catalogues are almost always rectangular\\footnote{The use of GALEX in WiggleZ \\citep{glaz07} is a rare counter-example.} and survey boundaries and volumes are easier to consider when using spherical coordinate derived edges. Packing a shape best described in spherical coordinates into a Cartesian defined region is a non trivial task, and many approaches have been used in redshift surveys. Such packing problems are of wider mathematical interest because no provably optimal and rapid technique has yet been discovered \\citep{megi84}. Instead every large survey tailors a tiling method in line with specific survey goals using a heuristic method. In this sense a heuristic method is one informed by knowledge of the problem at hand, the hope being the solution is not much worse than optimal. On top of the generic problem of efficient tile packing, spectroscopic surveys also have to contend with extremely non-uniform and complex selection functions within the tiles themselves. The major cause for the non-uniformity is object exclusion, either due to fibre collisions or slit overlaps.  In the case of 2dFGRS an approach close to hexagonal packing was used, where slight perturbations were made to a purely hexagonal grid of tile centres in order to better sample object densities. Since this survey was almost single pass (there was $\\sim 30\\%$ tile overlap), low completeness fields were not uncommon, an effect that was statistically adjusted for with observational weights. However, in the densest fields some targets will not have redshifts, and galaxy group assignments will not be as secure as in highly complete regions. The downside of such a regular approach is that all multi-fibre spectrographs will have structure or bias in their assignments and thus power will be added to (or removed from) certain frequencies in tangential modes of the power spectrum. The distribution of fibres is not only driven by the algorithm used to place them, but also the physical limitations of the instrument. Typically a fibre fed MOS is designed with fibres around the circumference in such a way that all fibres can reach the centre and few can reach locations at the edge, a scenario that makes radially-dependent targeting distributions inevitable. Even with the newest simulated annealing (SANN) algorithms available for AAOmega, radial assignment dependencies within each 2dF pointing exist \\citep{misz06}. It is obviously important to try to compensate for such biases in any work that is concerned with clustering and structure, such as Galaxy And Mass Assembly \\citep[GAMA,][]{driv09}, the latest large survey to use AAOmega on the AAT.  The spectroscopic element of SDSS \\citep{blan03} used a heuristic algorithm that attempted to find an acceptable solution of a perturbed uniform grid of tiles, much like 2dFGRS. The algorithm aimed to utilise 90\\% of the 600 available fibres on each tile, and similar to 2dFGRS the SDSS's median tile coverage for an object was 1 (both achieved a target density $\\sim$100 galaxies per square degree). Minimum fibre spacings are $55''$ for the SDSS spectrograph, larger than the $40''$ distance for 2dF, thus an obvious limitation of SDSS is the full targeting and unbiased analysis of close pairs (a key science objective for GAMA, discussed in detail below).  Of recent surveys, the VIMOS VLT Deep Survey (VVDS from here) utilised the simplest approach to tiling \\citep{bott05}. Effectively it placed tile centres on a fixed square grid with diagonal offsets used for the deeper component of the survey. Such an approach is possible when using VIMOS because of its mask-based grism spectrograph, giving it a square FOV better suited to tiling a square CCD photometric survey. The VVDS does not suffer from any radial selection bias, but due to the constraints imposed by slits cut into each mask it does possess complex selection effects such as the tendency to target a uniform spread of targets; highly clustered regions are hard to target since the slits necessarily avoid each other. Further complicating matters is a partially radial completeness bias, evident in the spectroscopic masks created for zCOSMOS \\citep{knob09}. Whilst an interesting survey to note, such a survey design is not trivial to create with any of the fibre based multi object spectrographs discussed due to their circular FOV, and the complex radial bias this introduces.  Simulated annealing solutions of the tiling problem have been utilised in large area surveys with large amounts of structure present, most notably by the 6dFGS \\citep{camp04}. Simulated annealing is a popular approach for many algorithmically insolvable problems and is, strictly speaking, a metaheuristic solution (i.e.\\ choices have to be made about the element to be optimised and also the method of optimising). In simple terms the user must pick something to be minimised (or maximised), such as the total number of objects not assigned to a fibre after tiling the whole survey region. The user must also give the SANN algorithm variables to perturb (most obviously the right ascension and declination of the tile centres), and a rate at which it `cools' towards a solution. Typically these perturbations become smaller as the solution improves, and eventually an acceptable set of tile positions should be found. Packing problems lend themselves well to SANN since they can be tuned to find acceptable solutions rapidly, but they are non-deterministic algorithms (unlike the other heuristic approaches discussed) and are neither provably optimal nor stable (i.e.\\ small variations to the problem to be optimised can produce radically different results). In the case of the 6dFGS, SANN is obviously much more effective than any sort of regular tiling because the projected target densities vary significantly and the survey area is large. The use of SANN reduces the number of sparsely populated fields and better samples overdensities where fields would be full.  Added to the complexities of these different approaches are the observational limitations for any survey as well as its scientific priorities. It will not be the case that all fields are equally observable in a large area survey (e.g.\\ varying rising and setting time as a function of RA), but in a sufficiently small area survey it will often be the case that all parts of the survey field are effectively as observable as each other. Also, the end point of the survey will often be an unknown (i.e.\\ weather dependent), so in many applications it is advantageous for the survey to be in a useable state as quickly as possible. With these extra considerations in mind, the philosophy that was applied to tiling GAMA was one where each tile would in some sense be the next most optimal tile, and every subsequent tile should make a significant impact towards achieving the GAMA survey requirements.  The GAMA redshift survey is one component of the multi-band GAMA survey project, and is the latest large survey to use the AAT's MOS facility. In this paper we explore the problems of tiling specifically for the GAMA survey, with the possibility of using the approaches discussed in future redshift surveys with characteristics in common with GAMA. In section 2 we outline the GAMA survey, and how the scientific goals for the project translate into survey requirements that our tiling algorithm must achieve. In section 3 we discuss in detail the different options to tiling that are appropriate for GAMA. In section 4 we apply the two most likely candidates for the tiling algorithm to the GAMA survey as it was left at the end of year 1, allowing quantitative judgements of the different approaches to be made. In section 5 we apply our chosen tiling algorithm to the data and present the state of the survey after year 2. Finally, conservative predictions are made for the state of the survey after year 3 observations based on tiling simulations. ",
        "conclusions": "This work has demonstrated that a greedy approach to tiling proves to be extremely successful in densely packed surveys such as GAMA. By aggressively targeting under-densities with each field used, high levels of spatial completeness should be a reality for each GAMA region by the end of the third year of observations. In the meantime we allow for a simple mechanism to feed redshift failures back in, and by prioritising highly clustered regions we obtain both a large number of close pairs, and guarantee we are not left with difficult pockets of galaxies in the final stages of the survey. Further still, by utilising every non-main survey fibre on deeper targets, we ensure efficient use of the 2dF instrument, and make a head-start on any future extended redshift surveys in the GAMA regions."
    },
    "0910/0910.0704_arXiv.txt": {
        "abstract": "Infrared spectroscopy of the \\ha{} emission lines of a sub-sample of 19 high-redshift (0.8 $< z <$ 2.3) Molonglo quasars, selected at 408 MHz, is presented. These emission lines are fitted with composite models of broad and narrow emission, which include combinations of classical broad-line regions of fast-moving gas clouds lying outside the quasar nucleus, and/or a theoretical model of emission from an optically-thick, flattened, rotating accretion disk, with velocity shifts allowed between the components. All bar one of the nineteen sources are found to have emission consistent with the presence of an optically-emitting accretion disk, with the exception appearing to display complex emission including at least three broad components. Ten of the quasars have strong Bayesian evidence for broad-line emission arising from an accretion disk together with a standard broad-line region, selected in preference to a model with two simple broad lines. Thus the best explanation for the complexity required to fit the broad \\ha{} lines in this sample is optical emission from an accretion disk in addition a region of fast-moving clouds. We derive estimates of the angle between the rotation axis of the accretion disk and the line of sight. Deprojecting radio sources on the assumption of jets emerging perpendicular to the accretion disk gives rough agreement with expectations of radio source models. The distribution in disk angles is broadly consistent with models in which a Doppler boosted core contributes to the chances of observing a source at low inclination to the line of sight, and in which the radio jets expand at constant speed up to a size of $\\sim 1$ Mpc. A weak correlation is found between the accretion disk angle and the logarithm of the low-frequency radio luminosity. This is direct, albeit tenuous, evidence for the receding torus model first suggested by \\citet{lawrence91} in which the opening angle of the torus widens with increasing radio luminosity. The highest accretion disk angle measured is 48$\\degree$, consistent with the opening angle predicted for radio-luminous sources. Velocity shifts of the broad \\ha{} components are analysed and the results found to be consistent with a two-component model comprising one single-peaked broad line emitted at the same redshift as the narrow lines, and emission from an accretion disk which appears to be preferentially redshifted with respect to the narrow lines for high-redshift sources and blueshifted relative to the narrow lines for low-redshift sources. An additional analysis is performed in which the disk emission is fixed at the redshift of the narrow-line region; although only two quasars show a robust change in fitted angle, the radio luminosity -- disk angle correlation falls sharply in probability, and so is strongly model dependent in this sample. ",
        "introduction": "\\subsection{Orientation effects} \\label{sec:orientation}  Optical emission spectra of Active Galactic Nuclei (AGN) matched in radio and optical luminosity are now firmly believed to be strongly influenced by orientation effects, with only small underlying differences in the sources themselves. There are two separate orientation-dependent effects which alter the optical spectra of AGN.  The Type 1/Type 2 classification of AGN is made according to the presence of broad emission lines. Type 2 AGN possess only narrow emission lines, $\\lesssim 2000$ \\kms{} \\citep[e.g.][]{peterson97} and weak non-stellar continuum emission. Type 1 AGN have broad emission lines of $\\sim 2000$ -- $20000$ \\kms, and strong non-stellar continuum emission, in addition to narrow lines similar to the Type 2 sources. An explanation for this disparity grew from the discoveries of broad emission lines seen in polarised light from Type 2 Seyfert sources \\citep[e.g.][]{antonucci85}, which suggested orientation-dependent obscuration caused by an intervening screen of matter, such as a dusty molecular torus \\citep{krolik86}. It has recently become clear that obscuration by dust in starbursting galaxies can also be responsible for concealing Type 2 AGN \\citep[e.g.][]{martinez05}.  The second orientation-dependent effect arises from the relativistic motion of the plasma in the radio jets. \\citet{scheuer79} first suggested that viewing a radio source with opposing relativistic jets would cause a large contrast in the luminosities of the approaching and receding jets, and that objects with jet axes close to the line of sight would be seen more often due to Doppler boosting of the core. \\citet{orr82} made the connection between Doppler boosting of the core, and a measure of quasar orientation from the core-to-lobe radio flux density ratio; this allowed them to unify the ``core-dominated'' quasars with flat optical spectra, viewed at angles close to the line of sight, with the ``steep-spectrum'' or radio-lobe-dominated quasars viewed at larger angles. \\citet{wills86} linked the radio properties to the optical properties by discovering an anticorrelation between the core-to-lobe radio flux ratio and the width of the \\hb{} line, interpreting this connection as the result of beaming of radio emission from a jet emerging along the rotation axis of an accretion disk; the broad \\hb{} lines arise from the accretion disk with a width correlated with the angle between the line of sight and the disk axis.  The two optical schemes were reconciled by \\citet{barthel89}, who gave a consistent picture in which FRII narrow line radio galaxies, steep-spectrum quasars and flat-spectrum quasars are all drawn from the same parent population, but viewed at decreasing angles to the line of sight. A review of these so-called unified schemes for AGN can be found in \\citet{urry95} or \\citet{antonucci93}.  A solid understanding of how quasar emission lines arise, and how they are affected by the AGN environment along different sight lines, is not only an interesting study in terms of quasar structure, but is also vital in order to disentangle orientation effects from large-scale AGN surveys which enable the study of cosmic evolution.  \\subsection{Accretion disks} \\label{sec:disks}  There is a growing body of evidence that AGN are powered by accretion of gas and dust onto supermassive black holes, and this is now the accepted paradigm. As the black hole feeds on the surrounding material, it is expected that this will form an accretion disk of infalling matter \\citep{shakura73}. The current theory is that accretion disks have two parts: a puffed-up, X-ray-emitting inner disk, and a flattened, outer disk which gives rise to broad optical emission lines.  \\citet{collin80} first suggested that a thickened inner accretion disk, shielding and reprocessing the hard X-ray emission from the black hole, gives rise to the optical low-ionisation \\feii{} emission seen in Type 1 Seyferts, from the atmosphere above the geometrically-thin outer part of the disk. \\citet{rees82} postulated that a geometrically- and optically-thick ion-supported torus surrounds the supermassive black hole at the centre of a radio galaxy, collimating the emerging radio jets. \\citet{tanaka95} observed a broad, asymmetric iron K$\\alpha$ line consistent with emission from a disk of this description, situated between $\\sim 3$ -- $10 \\rg$ (where $\\rg = G M / \\mathrm{c}^2$ is the gravitational radius) from the nucleus of the AGN.  \\citet{filippenko88} reviewed the different lines of evidence from optical and UV data that flattened, extended accretion disks fuel the central black holes of galaxies. For example, \\citet{baldwin77} recorded the anticorrelation of the UV continuum luminosity with the equivalent widths of broad \\civ{} emission (the ``Baldwin Effect''). Both this observation and the ``big blue bump'' of excess UV continuum emission found by \\citet{malkan82} may be explained by emission from an optically-thick, geometrically-thin accretion disk \\citep{netzer85}. The strongest direct evidence for this scenario is the presence of double-peaked, low-ionisation optical emission lines seen in some AGN, which arise from the Doppler effect acting on the emission from rotating material in the outer accretion disk (e.g. \\citet{chen89}, \\citet{perez88}).  The thin, optically-emitting disk must be illuminated by some mechanism. The photoionising flux may originate either from the central non-thermal source, or from the inner, X-ray-emitting disk \\citep{collin80}; and the radiation may illuminate the outer disk directly, or be scattered from a highly-ionised diffuse medium above the outer disk \\citep{chen89b}.  Optical double-peaked lines have to date only been found in a relatively low percentage of radio-loud AGN ($\\sim$ 10$\\%$, see \\citet{eracleous94}). \\citet{strateva03} discovered that radio-quiet quasars also emit double-peaked lines, although these appear to be rarer still: double-peaked emission was seen in $\\sim$ 3$\\%$ of the SDSS AGN, including both radio-loud and radio-quiet sources. Double-peaked profiles are not unique to Balmer lines: \\citet{strateva03} found double-peaked \\mgii{} emission lines in some SDSS AGN.  It is not clear why the double-peaked line profiles should be rare, although there are several possibilities: the outer accretion disk may simply be obscured by broad-line-emitting clouds surrounding it; or if the outer accretion disk causes a wind, then the broad lines which arise from this wind are predicted to be single-peaked \\citep{murray97}. In any case, it should not be expected that the low-ionisation broad lines seen in an AGN originate solely from the rotating disk; single-peaked emission may be seen in addition to double-peaked profiles.  The accretion disk model used in this paper is taken from \\citet{chen89}, and consists of a thick, hot torus, whose outer edge may reach up to $100\\, \\rg$ from the black hole. Inverse Compton-scattered X-rays from this inner disk illuminate an optically-thick, flattened outer disk \\citep{halpern89}. The thin, circular disk, which produces the double-peaked emission lines, may extend up to $\\sim 10^5\\, \\rg$ from the central engine. The distinctive line profiles are caused by the rotation of the disk, splitting the emission into redshifted receding material and blueshifted approaching material; the blueshifted peak is of higher intensity than the redshifted peak, as a result of Doppler boosting.  The chosen model was necessarily simple, to limit the number of free parameters.  More complex disks, such as an elliptical disk, which may arise when a single star is disrupted near the black hole \\citep{gurzadyan79}, or a warped disk, thought to occur around rotating black holes \\citep{bachev99}, give rise to a wider range of line profiles, including double-peaked profiles with a redward peak of higher intensity than the blueward peak. \\citet{strateva03} found that assuming all low-ionisation broad-line emission comes from a disk, non-axisymmetric disks would be required in $\\sim 60\\%$ of their SDSS sample, while \\citet{eracleous03} found that in their sample of 106 radio-loud AGN, 20\\% have double-peaked \\ha{} profiles visible to the eye, of which $\\sim 40\\%$ require a model more complex than the circular Keplerian disk.  In this paper, a small, but close to complete, sample of radio-loud quasars are analysed to determine if their spectra include emission from circular, planar accretion disks, either as the sole component of broad optical emission, or in combination with a single-peaked broad line.  \\subsection{Velocity shifts} \\label{sec:velshifts}  It has been generally accepted that the narrow-line region (NLR) of an AGN falls near the systemic redshift. The NLR is extended, and appears to be reasonably independent of viewing angle effects, and so the narrowness of the lines constrains the velocity of this gas to be small. \\citet{heckman81} showed, for a handful of low-$z${} Seyferts and radio galaxies, that the narrow lines had small blueshifts of between $\\sim$ 0 -- 300 \\kms{} relative to the neutral hydrogen emission of the host galaxies. \\citet{vandenberk01} discovered that, for composite spectra created with several thousand Sloan Digital Sky Survey (SDSS) quasar spectra covering a wide redshift range ($0.04 \\lesssim z \\lesssim 4.8$), the narrow lines are blueshifted by small velocities ($\\lesssim 100$ \\kms) which correlate with ionisation potential.  The broad-line region (BLR) has a complex structure, with emission line redshifts which vary according to species, implying separate regions of gas (see Figure 6 of \\citet{collin80}). It is thought that the high-ionisation lines (HILs) arise from a compact, spherical region close to the central black hole, while the low-ionisation lines (LILs) are formed in a flattened structure further from the AGN centre, possibly an accretion disk \\citep{krolik91}.  \\citet{collin88} suggested that the HILs might be produced by shocks in an outflowing wind.  \\citet{gaskell82} demonstrated for a sample of flat-spectrum quasars with z $\\sim$ 0.2 -- 2.3 that the HILs, such as \\ciii, \\civ{} and \\nv{}, are blueshifted by $\\sim 600$ \\kms{} with respect to the LILs, which include \\mgii{}, \\oi{} and the Balmer lines. \\citet{wilkes84} confirmed this shift for high-$z${} quasars ($2 \\lesssim z \\lesssim 3$), finding a slightly higher range of shifts, up to $\\sim$ 1400 \\kms. Blueshifted HIL zones have also been observed by \\citet{espey89}, who found shifts of $\\sim$ 1000 \\kms{} in a small sample of $1.3 \\lesssim z \\lesssim 2.4$ sources, and \\citet{corbin90}, who found shifts exceeding 4000 \\kms{} for luminous, optically-selected sources with $z \\gtrsim 1$. \\citet{richards02} also confirm this trend from studies of $\\sim 800$ quasars with $1.5 \\lesssim z \\lesssim 2.2$ from the SDSS, though they find a wide distribution of velocity shifts of the high-ionisation \\civ{} with respect to the low-ionisation \\mgii, ranging from redshifts of $\\sim$ 500 \\kms{} to blueshifts of over 2000 \\kms, and they take pains to point out that they do not believe it is a line shift so much as a lack of flux in the red wing of the line.  \\citet{mcintosh99} found that the HIL zone is at the same redshift as the narrow lines in low-$z${} sources, but for a sample of quasars with $2 \\lesssim z \\lesssim 2.5$, broad \\hb{} had a redshift of $\\sim$ 500 \\kms{} relative to \\oiii.  Although the LILs are typically considered to have low velocity shifts, there have been recorded instances of Balmer lines with very high redshifts, e.g. $\\sim$ 2100 \\kms{} for 3C277 \\citep{osterbrock76} and $\\sim$ 2600 \\kms{} for OQ208 \\citep{osterbrock79}.   \\section[]{Sample Selection} \\label{sec:selection}  A sub-sample of 19 quasars was defined from the Molonglo Quasar Sample \\citep{molonglo3}. The Molonglo sample of radio sources was selected from the Molonglo Reference Catalog (MRC) \\citep{large81}, a 408 MHz survey conducted with the Molonglo Synthesis Telescope which is 99.9\\% complete at 1 Jy. The radio sources were then identified from VLA 1-arcsecond resolution radio images, optical imaging and spectroscopy, the quasars \\citep{molonglo3} being distinguished from the radio galaxies \\citep{molonglo1} by the presence of broad optical emission lines.  The Molonglo quasar sample includes all quasars with flux densities \\mbox{$S_{408 \\mathrm{MHz}} > 0.95$ Jy} in a $10\\degree$-wide strip in the Southern sky, \\mbox{$-20\\degree > \\delta > -30\\degree$}, excluding sources with low Galactic latitude \\mbox{($\\mid b \\mid > 20\\degree$)} and also a strip in Galactic R.A. (details in \\citet{molonglo3}) to define a sample of manageable size. It should be noted that as the Molonglo sample was selected at the mid-range frequency of 408 MHz, there are likely to be some sources included in the sample by virtue of their strong radio core emission, and are therefore inclined at small angles to the line of sight; this builds an orientation bias into this survey.  The quasar sub-sample was selected on the basis of four criteria: observability during the relevant time period; redshift such that \\ha{} and \\hb{} emission falls in wavelength windows corresponding to regions of high atmospheric transparency ($0.8 < z < 1.0$, $1.5 < z < 1.65$, $2.2 < z < 2.3$); sufficient J-band brightness to be observable in a reasonable integration time; and exclusion of the RA range 14h -- 03h in accordance with scheduling constraints. The J-band magnitude limit is the only factor which adds a significant bias in the sub-sample selection. The objects were selected to be brighter than \\mbox{J $\\sim 18.5$}, and this excluded one source from the sample, MRC0418-288. This source is likely to be reddened or dusty, which means that it has a higher chance of being inclined at a large angle to the line of sight. MRC1256-243 is an extra core-dominated source which was not observed due to scheduling limitations; this source is likely to have been boosted into the sample by virtue of its strong core emission in any case.  Throughout the paper, a cosmology of $H_0 = 70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\mathrm{M}} = 0.3$ and $\\Omega_{\\Lambda} = 0.7$ is assumed. ",
        "conclusions": "\\label{sec:conclusions}  All bar one of the 19 quasars are best fitted with an emission model which includes more than one component of broad-line emission, in both analyses; there is overwhelming evidence that in this sub-sample of quasars, a simple one-component broad-line emission model is insufficient to describe the physical processes at work. There is strong Bayesian evidence that ten out of nineteen quasars from the analysis in which a shift is permitted between the BLR and NLR possess accretion disks, and all but one are consistent with the emission model of a circular Keplerian accretion disk in addition to a single-peaked, symmetric emission line arising from a separate BLR. The one quasar for which this does not apply appears to have very complex broad emission, consisting of at least three components. If the model does not allow for a shift between the BLR and NLR, strong evidence for disks is seen in eight out of the nineteen sources, and all but three are consistent with the presence of a circular accretion disk.  \\citet{eracleous94}, \\citet{eracleous03} and \\citet{strateva03} found that only 3 -- 20$\\%$ of AGN have obvious double-peaked lines (see Section \\ref{sec:disks}), in contrast to the $\\sim$ 84 -- 95\\% of this sample of quasars which were found to be consistent with the presence of a thin, optically-emitting accretion disk, despite only a few of the line profiles appearing obviously double-peaked to the eye. It may be that accretion disk emission has been seen in so few AGN because single-peaked broad emission lines, arising from fast-moving clouds located outside the thin disk or from a wind originating from the disk itself, have a tendency to swamp the disk emission. In these cases, a complex model and high signal-to-noise ratio spectra are required in order to retrieve the accretion disk contribution to the emission. It is therefore essential to extend this analysis to a larger sample of low-radio-frequency-selected quasars; optically-selected quasars are predicted to cover a narrower range in disk angles, as even lightly reddened quasars (at $\\sim 45 \\degree$) are much more likely to fall through the magnitude limit of a survey, and the optical emission may be angle dependent.  There is a possible correlation of the fitted disk angles with the projected source size, in agreement with the expected projection effects if the accretion disk is perpendicular to the jet. There are two CSS sources which do not appear to lie on the same relation as the other sources: MRC0222-224 seems to be an intrinsically smaller and more reddened source than the others, while MRC1114-220 may either be an outlier, or have a precessing or misaligned disk. When deprojected using the best-fit disk angles, the source sizes are consistent with a model in which the accretion disk is perpendicular to the radio jets, and the heads of the radio jets expand at approximately constant speed up to a size $\\sim 1$ Mpc; the paucity of sources larger than this limit can be explained by the average duty cycle of these AGN.  The distribution of fitted disk angles is consistent with the calculation of expected jet angles for Doppler-boosted jets with Lorentz factors of $\\sim$ 20, if the accretion disk rotation axis and radio jets are assumed to be coincident. The calculated jet angle distribution may match the disk angle distribution better if a more physically reasonable model was used for the jet, with a core of higher Lorentz factor than the surrounding jet sheath, as opposed to this simple single Lorentz factor model.  There is a weak correlation ($> 1.5 \\sigma$ significance by Kendall's $\\tau$ test) between the low-frequency radio luminosity and the sine of the angle between the accretion disk axis and the line of sight. This is predicted by the receding torus model of \\citet{lawrence91} if the accretion disks are perpendicular to the radio jets, such that the disk angles match the jet angles. In this model, the dust is sublimated by the AGN nucleus, leading to larger opening angles of the obscuring dusty torus for more radio-luminous sources, and hence quasars are seen up to larger jet angles. The largest accretion disk angle measured is 48$\\degree$, which is consistent with the opening angle predicted for powerful FRII radio sources. A larger sample size is needed to confirm the radio luminosity -- disk angle correlation, but this is vital, as it would yield the first direct test of the receding torus model. It should be noted that when the analysis was performed again with the models in which the accretion disk was fixed at the redshift of the NLR, this result disappeared, and so the radio luminosity -- disk angle correlation is strongly dependent on the model used.  A study of the velocity shifts between the line components revealed the single-peaked broad lines to be formed at similar redshifts to the narrow lines. The disk emission tended to be redshifted with respect to these lines for high-$z${} sources and blueshifted for lower-$z$ sources. Since the accretion disk is expected to be at the systemic velocity, this would imply that the single-peaked broad lines and the narrow lines are formed in outflows for more luminous, high-$z${} sources, and in infalling clouds for closer, less luminous sources. Such a scenario could arise if powerful winds were causing outflows of gas from the high-$z$, luminous sources, while these winds have stalled in the less-luminous quasars seen at $z < 1$, resulting in infall of the gas back towards the galaxy nucleus.  The results of this paper are all dependent on the validity of the disk model used in the Bayesian fitting. This model, taken from \\citet{chen89}, describes a circular, Keplerian disk. The illuminating flux from the central black hole was modelled as falling off with a radial exponent of 3. In order to investigate the effects of range of permutations to this basic model, such as a variety of disk emissivities and warped or elliptical disks, greater spectral resolution, a larger quasar sample size, and in particular, multi-epoch data to study the variability of the emission line profiles are required."
    },
    "0910/0910.2701_arXiv.txt": {
        "abstract": "The ``Mouse'' (PWN~G359.23$-$0.82) is a spectacular bow shock pulsar wind nebula, powered by the radio pulsar J1747$-$2958. The pulsar and its nebula are presumed to have a high space velocity, but their proper motions have not been directly measured. Here we present 8.5 GHz interferometric observations of the Mouse nebula with the Very Large Array, spanning a time baseline of 12 yr. We measure eastward proper motion for PWN~G359.23$-$0.82 (and hence indirectly for PSR~J1747$-$2958) of $12.9\\pm1.8$~mas~yr$^{-1}$, which at an assumed distance of 5~kpc corresponds to a transverse space velocity of $306\\pm43$~km~s$^{-1}$. Considering pressure balance at the apex of the bow shock, we calculate an in situ hydrogen number density of approximately $1.0_{-0.2}^{+0.4}$~cm$^{-3}$ for the interstellar medium through which the system is traveling. A lower age limit for PSR~J1747$-$2958 of $163_{-20}^{+28}$~kyr is calculated by considering its potential birth site. The large discrepancy with the pulsar's spin-down age of 25~kyr is possibly explained by surface dipole magnetic field growth on a timescale $\\approx$15~kyr, suggesting possible future evolution of PSR~J1747$-$2958 to a different class of neutron star. We also argue that the adjacent supernova remnant G359.1$-$0.5 is not physically associated with the Mouse system but is rather an unrelated object along the line of sight. ",
        "introduction": "The evolution of neutron stars and the potential relationships between some of their observed classes remain outstanding problems in astrophysics. Proper motion studies of neutron stars can provide independent age estimates with which to shed light on these questions. In particular, the well defined geometry of bow shock pulsar wind nebulae (PWNe; \\citeauthor{gaensler:3} \\citeyear{gaensler:3}), where the relativistic wind from a high-velocity pulsar is confined by ram pressure, can be used as a probe to aid in the understanding of both neutron star evolution and the properties of the local medium through which these stars travel.  The ``Mouse'' (PWN~G359.23$-$0.82), a non-thermal radio nebula, was discovered as part of a radio continuum survey of the Galactic center region \\citep{yusef}, and was suggested to be powered by a young pulsar following X-ray detection \\citep{predehl}. It is now recognized as a bow shock PWN moving supersonically through the interstellar medium (ISM; \\citeauthor{gaensler:5} \\citeyear{gaensler:5}). Its axially symmetric morphology, shown in Figure \\ref{fig:yusef20cm}, consists of a compact ``head'', a fainter ``body'' extending for $\\sim$10$^\\prime$$^\\prime$, and a long ``tail'' that extends westward behind the Mouse for $\\sim$40$^\\prime$$^\\prime$ and $\\sim$12$^\\prime$ at X-ray and radio wavelengths respectively \\citep{gaensler:5,mori}. The cometary tail appears to indicate motion away from a nearby supernova remnant (SNR), G359.1$-$0.5 \\citep{yusef}.  \\begin{figure*}[t] \\setlength{\\abovecaptionskip}{-7pt} \\begin{center} \\includegraphics[trim = 0mm 0mm 0mm 0mm, clip, angle=-90, width=13cm]{f1.ps} \\end{center} \\caption{VLA image of the Mouse (PWN~G359.23$-$0.82) at 1.4~GHz with a resolution of 12\\farcs8$\\times$8\\farcs4 (reproduced from \\citeauthor{gaensler:5} \\citeyear{gaensler:5}). The brightness scale is logarithmic, ranging between $-$2.0 and $+$87.6~mJy~beam$^{-1}$ as indicated by the scale bar to the right of the image. The eastern rim of SNR~G359.1$-$0.5 is faintly visible west of $\\sim$RA~17$^{\\mbox{\\scriptsize{h}}}$46$^{\\mbox{\\scriptsize{m}}}$25$^{\\mbox{\\scriptsize{s}}}$.} \\label{fig:yusef20cm} \\end{figure*}  A radio pulsar, J1747$-$2958, has been discovered within the ``head'' of the Mouse \\citep{camilo:1}. PSR~J1747$-$2958 has a spin period $P=98.8$ ms and period derivative $\\dot{P}=6.1\\times10^{-14}$, implying a spin-down luminosity $\\dot{E}=2.5 \\times 10^{36}$~ergs~s$^{-1}$, surface dipole magnetic field strength $B=2.5\\times10^{12}$~G, and characteristic age $\\tau_{c} \\equiv P/ 2\\dot{P}=25$~kyr (\\citeauthor{camilo:1} \\citeyear{camilo:1}; see also updated timing data from \\citeauthor{gaensler:5} \\citeyear{gaensler:5}). The distance to the pulsar is {\\footnotesize $\\gtrsim$}4~kpc from X-ray absorption \\citep{gaensler:5}, and {\\footnotesize $\\lesssim$}5.5 kpc~from HI absorption \\citep{uchida}. Here we assume that the system lies at a distance of $d=5d_{5}$~kpc, where $d_{5}=1\\pm0.2$ ($1\\sigma$).  Given such a small characteristic age, it is natural to ask where PSR~J1747$-$2958 was born and to try and find an associated SNR. While it is possible that no shell-type SNR is visible, such as with the Crab pulsar \\citep{sankrit} and other young pulsars \\citep{braun}, an association with the adjacent SNR~G359.1$-$0.5 appears plausible. This remnant was initially suggested to be an unrelated background object near the Galactic center \\citep{uchida}. However, it is now believed that the two may be located at roughly the same distance (\\citeauthor{yusef3} \\citeyear{yusef3}, and references therein). By determining a proper motion for PSR~J1747$-$2958, this association can be subjected to further scrutiny (for example, see analysis of PSR~B1757$-$24, PWN~G5.27$-$0.90 and SNR~G5.4$-$1.2; \\citeauthor{blazek} \\citeyear{blazek}; \\citeauthor{zeiger} \\citeyear{zeiger}).  As PSR~J1747$-$2958 is a very faint radio source, it is difficult to measure its proper motion interferometrically. It is also difficult to use pulsar timing to measure its proper motion due to timing noise and its location near the ecliptic plane \\citep{camilo:1}. To circumvent these issues, in this paper we investigate dual-epoch high-resolution radio observations of the Mouse nebula, spanning 12 years from 1993 to 2005, with the intention of indirectly inferring the motion of PSR~J1747$-$2958 through the motion of its bow shock PWN. In \\S~\\ref{SectionObservations} we present these observations. In \\S~\\ref{SectionAnalysis} we present our analysis and measurement of proper motion using derivative images of PWN~G359.23$-$0.82. In \\S~\\ref{SectionDiscussion} we use our measurement to determine an in situ hydrogen number density for the local ISM, to resolve the question of a possible association with SNR~G359.1$-$0.5, and to investigate the age and possible future evolution of PSR~J1747$-$2958. We summarize our conclusions in \\S~\\ref{SectionConclusions}. ",
        "conclusions": "We have investigated two epochs of interferometric data from the VLA spanning 12 years to indirectly infer a proper motion for the radio pulsar J1747$-$2958 through observation of its bow shock PWN~G359.23$-$0.82. Derivative images were used to highlight regions of rapid spatial variation in flux density within the original images, corresponding to the vicinity of the forward termination shock, thereby acting as a proxy for the motion of the pulsar.  We measure an eastward proper motion for PWN~G359.23$-$0.82 of $\\mu=12.9\\pm1.8$~mas~yr$^{-1}$, and interpret this value as an upper limit on the motion of PSR~J1747$-$2958. At this angular velocity, we argue that PSR~J1747$-$2958 is moving too slowly to be physically associated with the relatively young adjacent SNR~G359.1$-$0.5, independent of distance estimates to either object or of inclination effects.  At a distance $d=5d_{5}$~kpc, the proper motion corresponds to a projected velocity of $V_{\\mbox{\\tiny{PSR,$\\perp$}}}$~{\\footnotesize $\\leq$}~$\\left(306\\pm43\\right)d_{5}$~km~s$^{-1}$, which is consistent with the projected velocity distribution for young pulsars. Combining the time taken for PSR~J1747$-$2958 to traverse its smooth $\\sim$12$^\\prime$ radio tail with the time to escape a typical SNR, we calculate a lower age limit for PSR~J1747$-$2958 of $t_{\\mbox{\\tiny{total}}}$~{\\footnotesize $\\geq$}~$163_{-20}^{+28}$~kyr (68\\% confidence).  The lower age limit $t_{\\mbox{\\tiny{total}}}$ exceeds the characteristic age of PSR~J1747$-$2958 by more than a factor of 6, arguably providing the most robust evidence to date that some pulsars may be much older than their characteristic age. This age discrepancy for PSR~J1747$-$2958 suggests that the pulsar's spin rate is slowing with an estimated braking index $n$~{\\footnotesize $\\lesssim$}~$1.3$ and that its magnetic field is growing on a timescale $\\approx$15~kyr. Such potential for magnetic field growth in PSR~J1747$-$2958, in combination with other neutron stars that transcend their archetypal categories such as PSR~J1718$-$3718, a radio pulsar with a magnetar-strength magnetic field that does not exhibit magnetar-like emission \\citep{kaspi}, PSR~J1846$-$0258, a rotation-powered pulsar that exhibits magnetar-like behaviour \\citep{gavril,archibald}, and magnetars such as 1E 1547.0$-$5408 that exhibit radio emission (\\citeauthor{camilo:2} \\citeyear{camilo:2}, and references therein), supports the notion that there may be evolutionary links between the rotation-powered and magnetar classes of neutron stars. However, such a conclusion may be difficult to reconcile with evidence suggesting that magnetars are derived from more massive progenitors than normal pulsars (e.g., \\citeauthor{gaens:8} \\citeyear{gaens:8}; \\citeauthor{muno} \\citeyear{muno}). If the massive progenitor hypothesis is correct, then this raises further questioning of whether, like the magnetars, there is anything special about the progenitor properties of neutron stars such as PSR~J1747$-$2958, or whether all rotation-powered pulsars exhibit similar magnetic field growth or even magnetar-like phases in their lifetimes.  To constrain the motion of PSR~J1747$-$2958 further, future observational epochs are desirable. It may be possible to better constrain the motion and distance to this pulsar by interferometric astrometry with the next generation of sensitive radio telescopes (e.g., \\citeauthor{2004NewAR..48.1413C} \\citeyear{2004NewAR..48.1413C}). High time resolution X-ray observations may also be useful to detect any magnetar-like behaviour from this rotation-powered radio pulsar. In general, more neutron star discoveries, as well as measured or inferred braking indices, may allow for a better understanding of possible neutron star evolution."
    },
    "0910/0910.0368_arXiv.txt": {
        "abstract": "The role of null-point reconnection in a 3D numerical MHD model of solar emerging flux is investigated. The model consists of a twisted magnetic flux tube rising through a stratified convection zone and atmosphere to interact and reconnect with a horizontal overlying magnetic field in the atmosphere. Null points appear as the reconnection begins and persist throughout the rest of the emergence, where they can be found mostly in the model photosphere and transition region, forming two loose clusters on either side of the emerging flux tube. Up to 26 nulls are present at any one time, and tracking in time shows that there is a total of 305 overall, despite the initial simplicity of the magnetic field configuration. We find evidence for the reality of the nulls in terms of their methods of creation and destruction, their balance of signs, their long lifetimes, and their geometrical stability. We then show that due to the low parallel electric fields associated with the nulls, null-point reconnection is not the main type of magnetic reconnection involved in the interaction of the newly emerged flux with the overlying field. However, the large number of nulls implies that the topological structure of the magnetic field must be very complex and the importance of reconnection along separators or separatrix surfaces for flux emergence cannot be ruled out. ",
        "introduction": "\\label{sec:intro}  The continual injection of new magnetic flux, energy, and helicity into the Sun's atmosphere from the convection zone below is fundamentally responsible for much of the activity observed in the solar atmosphere. As the newly emerged flux interacts with the pre-existing magnetic field above, X-ray jets, bright points, or active-region loop systems may be formed (depending on the spatial scale of the emergence), thus these features are a result of magnetic reconnection and so are the realisation of the heated corona.  \\subsection{MHD modelling of flux emergence}  Numerical simulations are a powerful tool for modelling and understanding such complex inherently 3D systems. There is a large body of literature using magnetohydrodynamic (MHD) codes to model flux emergence events. Magnetic flux is created by the solar dynamo in the convective overshoot layer at the base of the convection zone, known as the tachocline \\citep{1992A&A...265..106S}. Instabilities cause tubes of flux to become buoyant and rise through the convection zone towards the solar surface. During the initial rise phase, the thin-flux-tube approximation has been used to predict active-region emergence latitudes and tilt angles that are consistent with observations \\citep{1994A&A...281L..69S,1995ApJ...441..886C,2006MNRAS.369.1703H}. However, as the flux tubes approach the photosphere they expand significantly, so the thin-flux-tube approximation is no longer valid, as the flux tubes are likely to become fragmented or at least severely distorted, and full 3D MHD models are required. The magnetic fieldlines must have enough twist to prevent the interactions with the surrounding medium from fragmenting the flux tube \\citep{1996ApJ...464..999L}, although the critical degree of twist required depends on the apex curvature of the loop and can be significantly less in 3D than the 2D limit implies \\citep{2000ApJ...540..548A}.  The emergence in three dimensions of a single flux tube or sheet into an initially field-free atmosphere was first studied by \\citet{1992PASJ...44..167M}, who obtained many basic features of the emerging field, such as draining of plasma down the fieldlines, expansion of loops into the corona, and formation of shock waves at the loop footpoints. Later work on this topic by several authors, all using a stably stratified convection zone, a low temperature photosphere and a high temperature (but field-free) corona, went into greater detail \\citep[see ][]{2005ApJ...635.1299A}. \\citet{2003ApJ...582..475A}, using an anelastic MHD model of the convection zone coupled to a fully compressible code for the atmosphere, found that the newly-emerged coronal fieldlines formed sigmoidal structures whose chirality depended on where in the atmosphere the emitting plasma was assumed to be located. \\citet{2003ApJ...586..630M} showed that the initially bipolar photospheric magnetic field structure can develop into a quadrupolar structure, and developed the classification of emerged fieldlines as either expanding or undulating. Finally, \\citet{2004ApJ...610..588M} found that the natural shearing motions that occur as the field expands into a pressure-stratified atmosphere tend to add axial flux and decrease the twist of the fieldlines.  However, the real corona contains a pre-existing magnetic field, and interesting results have been obtained by including an overlying field in flux emergence models. In 2.5D, \\citet{1996PASJ...48..353Y} experimented with vertical, oblique, and several orientations of horizontal background coronal magnetic fields. All the configurations lead to a current sheet forming over the emerging loops, and the production of magnetic islands and jets via the tearing instability. However, as \\citet{2003ApJ...586..630M} have pointed out, full 3D models are required to adequately reproduce the draining of plasma down the magnetic fieldlines. Such simulations have been carried out by, for example, \\citet{2004A&A...426.1047A,2005ApJ...635.1299A}, who allowed a twisted flux tube to emerge into a uniform horizontal background magnetic field. They found that in such a 3D setup, magnetic reconnection occurred at multiple sites along fieldlines, rather than at a single isolated magnetic null as is the norm in 2D. This model reproduced fan-like arcades similar to those observed by TRACE, and also reconnection jets whose temperatures and velocities were a good fit with observations.  \\citet{2004ApJ...609.1123F} also emerged a 3D twisted magnetic flux tube into a pre-existing coronal field, although here the coronal field was a potential arcade. Their model does not include a convection zone; the flux tube is `injected' into the atmosphere by controlling the value of the electric field ($\\mathbf{v} \\times \\mathbf{B}$) at the lower boundary, thus specifying the electric current and hence the Lorentz force. This allows them to achieve full emergence of the tube into the atmosphere, by analytically defining the shape of the field, and the velocity with which it passes through the boundary. Although the ideal MHD equations were used, magnetic reconnection could still take place at current sheets due to numerical diffusion in the presence of strong gradients in the magnetic field. Once a certain critical amount of twist had emerged, the flux tube underwent the kink instability and began to accelerate upwards and twist around, producing a sigmoidal current sheet beneath. This evolution was found to be strongly dependent on the relative orientations of the emerging flux tube and the pre-existing coronal magnetic field: if the overlying field was oppositely-directed to the axial component of the flux tube, reconnection started straightaway and destroyed the flux tube as it rose into the atmosphere. The effect of the orientation of the overlying field was investigated by \\citet{2007ApJ...666..516G}, using the same model as the Archontis papers above, but with different orientations of the plane-parallel coronal field relative to the emerging flux tube. They found that most reconnection took place when the two flux systems were close to anti-parallel, and hardly any when they were close to parallel, but, in all cases, the height reached by the emerging flux tube, as a function of time, was unchanged.  The models described so far all made certain simplifications about the behaviour of the convection zone, since their main aim was to investigate the effects of emergence into the solar atmosphere. However, the most realistic models currently possible of magnetic flux rising through the convection zone include radiative transfer, convection, compressibility, and thermal conduction. For example, \\citet{2007A&A...467..703C} showed how a single initial rising flux tube influenced by convection leads to the emergence of many small flux bundles, which produce a characteristic (observed) dark lane appearing in the photospheric granulation pattern. However, since they have no atmosphere above the photosphere they could not study the effects of emergence on the atmosphere itself. \\citet{2008ApJ...679..871M} recently performed a simulation in which their corona was heated by the release of stress built up by convective motions in a complex magnetic field topology. The granulation cells became enlarged as an emerging flux tube passed through them, in agreement with observations.  Recent work by \\citet{2007A&A...470..709M} has shown that in fact such a high level of complexity in the subsurface field is not necessary for emergence models to give valid results in the solar atmosphere. They create a complex subsurface field by allowing two magnetic flux tubes to interact in a stably stratified convection zone before emerging, and show that the resulting emerged atmospheric magnetic field is very similar to the result for a simple single flux tube. This means that the results obtained from simpler models with stably-stratified convection zones are adequate when the aim is to study the resulting atmospheric magnetic fields.  \\subsection{The next step: magnetic topology}  As we have just seen, previous work on solar emerging flux has greatly improved our understanding of many of the complex processes involved. But there is one key issue that was not addressed in all the previous analyses; they do not account for the topology of the magnetic fields. Here, by magnetic topology, we mean specifically the location and evolution of topological features such as magnetic null points, spines, separatrix surfaces, and separators \\citep[see][for a good review]{2005LRSP....2....7L}. Plotting individual fieldlines that are not part of the topological skeleton, even if they are carefully selected, is not enough to ensure that all the information about the structure of the magnetic field, and the way it connects, is known.  Magnetic reconnection preferentially occurs at topological features including null points \\citep{2004GApFD..98..407P,2005GApFD..99...77P,2007JGRA..11203103P}, separatrix surfaces \\citep{2002ApJ...576..533P,2005ApJ...624.1057P}, and separator fieldlines \\citep{2005ApJ...624.1057P,2007RSPSA.463.1097H,2008ApJ...675.1656P}, as well as at their geometrical counterparts: quasi-separatrix layers \\citep{2006SoPh..238..347A,2006AdSpR..37.1269D,2007Sci...318.1588A,2007ApJ...660..863T} and quasi-separators \\citep[also known as hyperbolic flux tubes;][]{2003ApJ...582.1172T,2003ApJ...595..506G,2005A&A...444..961A,2006A&A...451.1101D,2006A&A...459..627D,2007A&A...473..615W}. This is because the hyperbolic magnetic topologies around such features tend to focus the electric current, and a strong electric field (parallel to the magnetic field) is associated with 3D magnetic reconnection. So knowing the topological structure of the field provides a good guide to the probable locations of reconnection sites in the system.  In this paper, we concentrate on determining the importance of reconnection at magnetic null points in our model magnetic field. The signature of 3D magnetic reconnection is the existence of a strong electric field parallel to the magnetic field \\citep{1988JGR....93.5547S}. We therefore study the parallel electric field near the nulls, to determine if reconnection takes place there. We also consider the current density, which may be parallel ($\\mathbf{j}_{\\parallel}$) or perpendicular ($\\mathbf{j}_{\\perp}$) to the spine of the null point, or a combination of both, and the consequences of this for the structure of the surrounding fieldlines are discussed in Section~\\ref{ssec:locnat}. 3D magnetic reconnection may take place in the presence or absence of magnetic null points \\citep{1988JGR....93.5547S,1988JGR....93.5559H,1995JGR...10023443P,1996JGR...101.7631D,1996RSPSA.354.2951P,2003PhPl...10.2717H,2006RSPSA.462.2877W}. However, in the flux emergence model that we study here, the change of magnetic topological structure is a central feature, and is cospatial with the location of magnetic reconnection. To have a changing magnetic topology, nulls are often involved, and hence we focus on reconnection at nulls in this paper. Reconnection at other topological features in our flux emergence experiment may also be significant, and this will form the subject matter for a future investigation.  So, the magnetic topology is important as it allows us to identify probable sites of magnetic reconnection. However, it has only been realised very recently that information about the topological structure is also required to make sense of the reconnection rate and energy release sites. Knowing the rate at which magnetic reconnection takes place is crucial to our understanding of how the magnetic field evolves, and recent work by \\citet{2008ApJ...675.1656P} has shown that knowledge of the magnetic topology is absolutely essential for a correct interpretation of the results of the model. This is because, in a complicated magnetic field containing fieldlines with many different connectivities, a phenomenon called \\emph{recursive reconnection} can take place. Recursive reconnection means that the same magnetic flux can be recycled many times through a repeating sequence of different magnetic connectivities. In a situation like this, knowledge of the topological structure of the magnetic field is vital to determine the global reconnection rate, which can in fact be much higher than would be determined via other methods.  In this paper, we take the first steps towards a full topological analysis of a flux emergence model. The dataset that we use is one of the numerical MHD models of \\citet{2007ApJ...666..516G}, which consists of a twisted buoyant magnetic flux tube rising through the upper layers of a stably-stratified convection zone and into an atmosphere with a horizontal plane-parallel pre-existing magnetic field. Now, all of the fieldlines in the topological skeleton of the magnetic field must either start or end at magnetic null points (there are no bald patches \\citep{1993A&A...276..564T} in our model as no fieldlines leave the closed boundaries of our box, only the periodic boundaries). We locate the null points in the magnetic field using the accurate new algorithm described by \\citet{2007PhPl...14h2107H}. This paper concentrates on the surprising nature of these magnetic null points: their number, type, distribution, and evolution. Section~\\ref{sec:model} describes the code, the model setup, and how it was analysed. Our results are reported in Section~\\ref{sec:results} and then we conclude with a discussion in Section~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  In this paper, we have analysed the properties of the magnetic null points present in a model solar flux emergence event, and considered the implications for magnetic reconnection.  The flux emergence model that we used was first described by \\citet{2007ApJ...666..516G}. In it, a twisted magnetic flux tube rises through a stratified solar convection zone and atmosphere to interact and reconnect with a horizontal overlying magnetic field. We detected a large number of magnetic null points within the resulting model magnetic field. Nulls begin to appear when the magnetic reconnection starts, and persist until the end of the entire model run. Up to 26 nulls are present at any one time.  The nulls were tracked through all the time-frames of the model, and we found that they number 305 in total. Our surprise at this unexpectedly-large number of nulls led us to speculate that some of them might simply be numerical artefacts. However, there is a lot of evidence to show that the vast majority of nulls are real, are not due to numerical artefacts, and that we can reliably identify and exclude those that are. The nulls were also classified by sign (positive or negative), and we found that the balance of signs is always in agreement with a conservation law derived from the 3D Euler equation (equation~\\ref{eq:euler}), once the spurious nulls are excluded. The lifetimes of the nulls were calculated, and the majority of the nulls were found to persist for many multiples of the mean timestep between frames. Finally, the stability of the nulls was also assessed by calculating how fast their spine eigenvectors were rotating. The vast majority of nulls were found to have remarkably steady spine directions, meaning that they are very stable.  The nulls are located throughout the model photosphere and transition region. They never reach as high as the bottom of the model corona. The reason for this is because in the vast majority of our model corona the magnetic field is very simple (an overlying horizontal magnetic field) and the emergence of new flux only creates complexity in small localised regions low down in the atmosphere. \\citet{2009SoPh..254...51L} have investigated the number of nulls that occur in the solar atmosphere and showed that indeed their number falls off quickly the further you get from the photosphere, as the complexity of the magnetic field decreases with height. The nulls in our experiment form just after the reconnection has started, in two clumps low down at the boundary between the emerging and the overlying magnetic flux. As the emerging region expands, the nulls are pushed out towards the side boundaries of the box, but they never reach or cross these boundaries during the model run.  A parallel electric field is a necessary requirement for reconnection, so we investigated its nature in the vicinity of the nulls. Regions of strong parallel electric field are found in the model at the peak of the emerging fieldlines, and also (later on) lower down in the atmosphere. Although many of the nulls have some associated parallel electric field, we found that this is weak compared to the strongest parallel electric fields elsewhere in the atmosphere, which are concentrated far from the nulls in the regions just described. Null-point reconnection is therefore ruled out as the predominant type of magnetic reconnection taking place during the emergence event. Clearly though, the reconnection is associated with either separators from these nulls or reconnection in the absence of nulls. We plan to investigate these possibilities more thoroughly in future.  \\begin{acks} The authors would like to thank Dr.\\ D.\\ Pontin for his insightful comments, and also the anonymous referee for his thoughtful suggestions which helped us to improve the paper. Computational time on the UKMHD Linux clusters in St Andrews (STFC and SRIF funded) is gratefully acknowledged. This work was supported by the European Commission through the SOLAIRE Network. RCM is financially supported by the University of St Andrews Solar Group STFC Rolling Grant.  \\end{acks}"
    },
    "0910/0910.3513.txt": {
        "abstract": "{  Accurate photometric CoRoT space observations of a secondary seismological target, \\hk, led to the discovery that this star is an astrophysically important double-lined eclipsing spectroscopic binary in an eccentric orbit ($e\\sim0.3$), unusual for its short 3\\fd65705 orbital period. The high eccentricity, coupled with the orientation of the binary orbit in space, explains the very unusual observed light curve with strongly unequal primary and secondary eclipses having the depth ratio of 1-to-100 in the CoRoT ``seismo\" passband. Without the high accuracy of the CoRoT photometry, the secondary eclipse, 1.5 mmag deep, would  have gone unnoticed. A spectroscopic follow-up program   provided  45 high dispersion spectra. The  analysis of the CoRoT light curve was performed with an adapted version of PHOEBE that supports CoRoT passbands. The final solution was obtained by simultaneous fitting of the light and the radial velocity curves. Individual star spectra were derived by  spectrum disentangling. The uncertainties of the fit were derived by bootstrap resampling and the solution uniqueness was tested by heuristic scanning.  The results provide a consistent picture of the system composed of two late B stars.  The Fourier analysis of the light curve fit residuals yields two components, with orbital frequency multiples and an amplitude of $\\sim$ 0.1 mmag, which are tentatively interpreted as tidally induced pulsations.  An extensive  comparison with theoretical models is carried out by means of the Levenberg-Marquardt minimization technique and the discrepancy between models  and the derived parameters is discussed.  The best fitting models yield a young system age of 125 million years which is consistent with the eccentric orbit and  synchronous component rotation at periastron.\\\\} ",
        "introduction": "\\object{HD 174884} (Corot 7758) was selected as a  seismology target for   the first  short run in the CoRoT ``center\"  direction (SRc1) because  -- on the basis of its B8~V spectral classification  -- it was considered  a good candidate to pinpoint the instability strip red edge of the Slowly Pulsating B stars \\citep[e.g.][]{wae98}. Its only available spectrum, obtained before launch with the Elodie spectrograph and available from the Gaudi database\\footnote{GAUDI (sdc.laeff.inta.es/gaudi/) is the data archive  of the ground-based asteroseismology programme of the CoRoT mission. The GAUDI system is maintained at LAEFF, which is part of the Space Science Division of INTA.}  \\citep{sol05} did not reveal the star's binary nature, which was instead immediately evident from  CoRoT photometry. The system was continuously observed  during 27 days. The light curve of HD~174884 is characterized by signatures of high eccentricity: the secondary minimum markedly displaced from phase 0.5 and a  bump, typical of close eccentric binaries, due to variable distortion of stellar surfaces along the orbit. The system is very interesting for several reasons: the high eccentricity is coupled with quite short a period ($P_{orb}=3 \\fd 66$), putting the system in an extreme location of the period - eccentricity diagram of B stars. This indicates young age or an inefficient circularization process. Secondly, the huge difference between the primary and  secondary minima, the latter detected only thanks to CoRoT accuracy, and the overall shape and accuracy of the light curve is a challenge for light curve fitting methods and a test of their adequacy to solve light curves of this unprecedented quality. Finally, only a thorough analysis of HD~174884 could yield information of the possible pulsational properties of the components. \\begin{figure*}[!ht] \\centering  \\includegraphics[width=17.4cm]{13311fg1.eps}  % compressed version \\caption{Upper panel: the complete CoRoT light curve of HD 174884 as normalized flux vs HJD. Lower panel: a blow-up, spanning approximately three periods, showing the tiny secondary minimum (first occurrence at reduced HJD of around 4218) } \\label{lc} \\end{figure*} In the following sections we present the results of the analysis of CoRoT photometry and of ground-based high dispersion spectroscopy. A critical analysis of parameter uncertainties and of the uniqueness of the light and radial velocity curve  solution is an essential part of the study.  The well constrained physical parameters of the binary components are then compared with  evolutionary models computed with  the CL\\'ES code \\citep[Code Li\\`{e}geois d' \\'{E}volution Stellaire,][]{scuetal08} with the aim of constraining the age and the structural parameters (chemical composition, overshooting). Finally the light curve fit residuals, with variations of a few $10^{-4}$ mag, are analyzed.   %__________________________________________________________________ ",
        "conclusions": "\\label{conclusions} This work has been carried out with two aims in mind:  to achieve a thorough description of an  interesting binary and to test  the performance of  current binary star modeling on data  of unprecedented accuracy, as those available for the \\corot{} seismo-field targets. \\hd lends itself to both purposes:   its peculiar light curve suggests an unusual system configuration, worth a detailed study, and it is as well quite stable over the seven observed orbital periods so that,  if possible systematic errors arising from  binary modeling were present, they would not be hidden by transient intrinsic phenomena. (The latter is indeed the case of the other fully analyzed binary seismo-target of \\corot\\ first runs, \\object{AU Mon} \\citep{desmetal09} whose study reveals the presence of variable circumbinary material).  Given the presence of a grazing eclipse, we were concerned about solution uniqueness, a problem which is sometimes overlooked, and we put a great effort in getting a sound estimate of parameter uncertainties, essential for sensible comparison with evolutionary stellar models.  We can summarize our main results as follows: \\begin{itemize} \\item{due to the high  accuracy of CoRoT data, we were able to unambiguously derive parameters, such as inclination and component sizes, which would otherwise be poorly constrained in a system with grazing eclipses. This fact, the excellent agreement between the results of the different  methods applied for the analysis of the spectra, and the extensive check of solution uniqueness allowed a sound estimation of the physical parameters of \\hd  and their uncertainties.} \\item{The comparison with stellar  models  is quite satisfactory but for the temperature difference between the components, which, for the best fitting models,  is typically a few hundred degrees higher than that from the light curve  solution.  We interpreted that as due to the comparison with non rotating stellar models. We could as well improve the agreement with models introducing differences in the physical processes acting in the components, such as overshooting, but Occam's razor arguments  favor  the simpler, though not exhaustive, explanation. Increasing the free parameters yields, in fact, only a marginal improvement. The comparison with theoretical models provided as well an estimation of the system age and some indication on the component chemical composition.} \\item{The dynamical properties of \\hd are  in very good agreement with the predictions of Zahn's theory of circularization and spin-orbit synchronization. The high value of the eccentricity for the period suggests that  resonance locking has not been at work in the system.} \\item{A few frequencies are clearly detected in the residuals of the light curve fit at multiples of the orbital frequency.  We tentatively interpret them as resonantly excited pulsations. This hypothesis is strengthened by the fact that many g-modes of high radial order and degree $\\ell=2$ exist in the  frequency range of  tidal frequencies. Among the very few detections of tidally induced pulsations in binaries,  \\hd is the only case with two well characterized  early type MS components.} \\end{itemize}  Finally, it is evident from  our results that when dealing with observation  accuracy of a few tenths of mmag  we are close to the level at which systematics from the model has a relevant impact on the light curve solution. This indicates that the models shall be updated, especially in view of the still higher accuracy curves which will be obtained by current and future space missions such as  Kepler \\citep{kepler}. Similarly, the physical properties derived by data of such accuracy should be compared with theoretical models  including  stellar rotation.  %__________________________________________________________________"
    },
    "0910/0910.1587_arXiv.txt": {
        "abstract": "Supermassive black holes (SMBHs) found in the centers of many galaxies are understood to play a fundamental, active role in the cosmological structure formation process. In hierarchical formation scenarios, SMBHs are expected to form binaries following the merger of their host galaxies. If these binaries do not coalesce before the merger with a third galaxy, the formation of a black hole triple system is possible. Numerical simulations of the dynamics of triples within galaxy cores exhibit phases of very high eccentricity (as high as $e \\sim 0.99$). During these phases, intense bursts of gravitational radiation can be emitted at orbital periapsis, which produces a gravitational wave signal at frequencies substantially higher than the orbital frequency. The likelihood of detection of these bursts with pulsar timing and the Laser Interferometer Space Antenna ({\\it LISA}) is estimated using several population models of SMBHs with masses $\\gtrsim 10^7~{\\rm M_\\odot}$. Assuming 10\\% or more of binaries are in triple systems, we find that up to a few dozen of these bursts will produce residuals $>1$ ns, within the sensitivity range of forthcoming pulsar timing arrays (PTAs). However, most of such bursts will be washed out in the underlying confusion noise produced by all the other 'standard' SMBH binaries emitting in the same frequency window. A detailed data analysis study would be required to assess resolvability of such sources. Implementing a basic resolvability criterion, we find that the chance of catching a resolvable burst at a one nanosecond precision level is $2-50$\\%, depending on the adopted SMBH evolution model. On the other hand, the probability of detecting bursts produced by massive binaries (masses $\\gtrsim 10^7\\msun$) with {\\it LISA} is negligible. ",
        "introduction": "\\label{intro} It is well established that most galaxies host supermassive black holes (SMBHs) in their centers \\citep{rich98}. In the past decade, compelling evidence of the correlation between the mass of the central SMBH and the bulge  velocity dispersion and luminosity has been collected ~\\citep{ferrarese00,gebhardt00,merritt01,tremaine02}, indicating a coevolutionary scenario for SMBHs and their hosts. On a cosmological scale, galaxy formation and evolution can be understood by semi-analytic modeling, where properties of the baryonic matter are followed in the evolving dark matter halos obtained from large-scale models of hierarchical gravitational structure formation. A simple model of galaxy and central SMBH evolution in which every merger of galaxies leads quickly to coalescence of their central black holes can quantitatively reproduce both the SMBH mass-bulge luminosity relation~\\citep{kauffmann00} and the SMBH mass-velocity dispersion relation~\\citep{haehnelt00}.  In this general picture, if both of the galaxies involved in a merger host a SMBH, then the formation of a SMBH binary is an inevitable stage of the merging process. Following the merger, the two black holes sink to the center of the merger remnant because of dynamical friction ~\\citep{begelman80}. When the mass (either in gas or stars) enclosed in their orbit is of the order of their own mass, they start to feel the gravitational pull of each other, forming a bound binary. The subsequent binary evolution is, however, still unclear. In order to coalesce, the binary must shed its binding energy and angular momentum; a dynamical process known in literature as `hardening'. A crucial point in assessing the fate of the binary is the efficiency with which it transfers energy and angular momentum to the surrounding gas and stars.  The case of SMBH binaries in stellar environments has received a lot of attention in the last decade. The system is usually modeled as a massive binary embedded in a stellar background with a given phase space distribution. The region of phase space containing stars that can interact with the SMBH binary in one orbital period is known as the loss cone~\\citep{frankr76,AS01,milos03}. As the binary evolves, it ejects stars on intersecting orbits via the so called `slingshot mechanism', causing a progressive emptying of the loss cone, which ultimately increases the hardening time scale. Without an efficient physical mechanism for repopulating the loss cone, the binary will never proceed to small separations where coalescence induced by gravitational radiation takes place within a Hubble time. This is known as the stalling or `last parsec' problem \\citep{milos01}.  In the last decade, several solutions to the stalling issue have been proposed. Axisymmetric or triaxial stellar distributions may significantly shorten the coalescence timescale \\citep{yu02,merritt04,bercziketal06}. This is bacause the presence of deviations from spherical symmetry can produce ``boxy'' orbits, as seen by \\cite{bercziketal06}.  These orbits produce centrophilic stellar orbits and, therefore, replenish the loss-cone.  However, more recent calculations by \\cite{AmaroSeoaneSantamaria09} of the outcome of the merger of two clusters initially in parabolic orbits \\citep{ASF06} have not been able to reproduce the rotation necessary to create the unstable bar structure. Other studies have invoked eccentricities of the binary to refill the loss cone, since this effect could alter the cross section for super-elastic scatterings (thus altering the state of the loss cone) and shorten the gap to the onset of gravitational radiation effects (e.g.:~\\citealt{hemsendorf02, aarseth03, bercziketal06,ASF06,ASMF09}). The presence of massive perturbers may also help replenishing the loss cone, boosting the binary hardening rate \\citep{perets07}. On the other hand, in smooth particle hydrodynamics simulations of SMBH binaries in gas-rich environments, efficient hardening induced by the tidal interaction between the binary and the gas medium has been observed, indicating a possible quick coalescence \\citep{escala05,dotti06}. However, current simulations do not have the resolution to follow  the binary fate down to the gravitational wave (GW) emission regime, and robust conclusions about its late inspiral and coalescence can not be drawn. In any case, very massive low redshift systems, which are the major focus of our study, are more likely to reside in massive gas poor galaxies and their dynamics is probably dominated by stellar interactions.  When scaled to very massive binaries (masses $>10^8\\msun$), the inferred coalescence timescales in a stellar dominated environment are of the order of few Gyrs, indicating that SMBH binaries may be relatively long living systems. If the typical timescale between two subsequent mergers is comparable the SMBH binary lifetime, then a third black hole may reach the nucleus when the binary is still in place, and the formation of SMBH triplets might be a common step in the galaxy formation process. Recent studies of galaxy pairs lead to the conclusion that $30-70$\\% of present day massive galaxies have undergone a major merger since redshift one \\citep{bell06,lin08}, where 'major' means with baryonic mass ratio of the two components larger than $1/3$ or $1/4$ (depending on the study), which is a quite conservative threshold. This means that, on average, all massive galaxies have experienced a merger event in the last ten billion years. Assuming uncorrelated events, and a typical binary lifetime of one billion years, then 10\\% of SMBH binaries may form a triplet. With increasing redshift (and decreasing masses), dynamical timescales become shorter and shorter, implying that triplets may have been more common in the high redshift Universe.  In this paper we focus on SMBH triplets, studying their dynamical evolution, GW emission, and detectability. Employing sophisticated three body scattering experiments calibrated on direct-summation {\\sc Nbody} simulations, we study the dynamical evolution of the system, finding surprisingly high eccentricities of the inner SMBH binary (up to $e>0.99$). Even though the triple interaction would possibly lead to an ejection of one or even all SMBHs \\citep{valtonen94}, most of the systems are long living ($\\sim10^9$ yrs, \\cite{HoffmanLoeb07}), and final coalescence is more common than ejection, confirming analytical results by \\cite{makino94}. We model at the leading quadrupole order \\citep{peters63} the bursts of gravitational radiation emitted in the highly eccentric phase, assessing detectability with future GW experiments. Adopting cosmologically and astrophysically motivated models for SMBH formation and evolution, we estimate reliable event rates.  In order to cover the low frequencies generated by the expected cosmological population of coalescing SMBH binaries \\citep[e.g.,][]{wl03,ses04,ses05,ses07} or plunges of compact objects such as stellar black holes on to supermassive ones \\citep[see e.g.][for a review and references therein]{Amaro-SeoaneEtAl07}, the space-born observatory {{\\it LISA}} \\citep{bender98} has been planned to be covering the range of frequencies of $\\sim10^{-4}-10^{-1}~{\\rm Hz}$. Moving to even lower frequencies, the Parkes Pulsar Timing Array \\citep[PPTA,~][]{manchester06,manchester07}, the European Pulsar Timing Array \\citep[EPTA,~][]{EPTA} and the North American Nanohertz Observatory for Gravitational Waves \\citep[NANOGrav,~][]{NANOGrav} are already collecting data and improving their sensitivity in the frequency range of $\\sim10^{-8}-10^{-6}$ Hz, and in the next decade the planned Square Kilometer Array \\citep[SKA,~][]{laz09} will provide a major leap in sensitivity.  Throughout this paper we consider only very massive systems, with total mass $\\sim 10^8\\msun$. Our goal is to investigate if the high frequency nature of eccentric bursts can provide information about systems which would otherwise emit outside the frequency windows of the planned GW experiments quoted above, by shifting wide (separation $\\gtrsim0.1$ pc) SMBH binaries into the PTA window or by boosting relatively massive (masses $>10^7\\msun$) systems into the {\\it LISA} domain. We note that the bursts analyzed here are different from the `bursts with memory', which arise during the actual coalescence of SMBH binaries and are discussed in~\\citet{pshirkov09} and~\\citet{vanhaasteren09}.     The structure of the paper is as follows. In Section \\ref{astro}, we describe our comprehensive study of the dynamics of triple systems and investigate the eccentricity evolution of the inner binary by using direct-summation $N-$body techniques and a statistical 3-body sample calibrated on the $N-$body results. In Section \\ref{gravwaves}, we model the GW signal produced by eccentric bursts and we introduce observable quantities for PTAs and {{\\it LISA}}. In Section 4 we construct detailed populations of emitting SMBH binaries and triplets, and we discuss our results in terms of signal observability and detection rates in Section 5. Lastly, we briefly summarize our results in Section 6. ",
        "conclusions": "\\label{conclusions}  We have addressed in this work three different points in the evolution of triplets of SMBHs in the Universe: The Astrodynamics of the system,  the potential GW signature and the detectability.  We have performed eight different direct-summation $N-$body  simulations, one including more than half a million of particles, to calibrate 1,000 3-body scattering experiments, which include post-Newtonian corrections, in order to have a statistical description of the system. Both numerical tools agree that the inner binary of SMBHs will go through a phase of extremely high eccentricity, which is the motivation for the rest of the work.  These three-body excitations of episodic high eccentricity  configurations of the close SMBH binary produce interesting GW bursts that may be detectable with forthcoming experiments such as PTAs and {{\\it LISA}}.  The extreme eccentricities of such bursts on one hand would leave a very distinctive signature, but on the other require the development of appropriate analysis techniques.  To compute likely event rates, we extracted catalogues of merging galaxies from the Millennium Run, and we populated them with SMBHs following the known MBH-bulges relations. We then estimated the fractions of triplets and their eccentricity distribution and we computed the induced signals in both PTAs and the {{\\it LISA}} detector.  We found that, depending on the details of the SMBH population model, if the fraction of triplets is $\\ge 0.1$, few to a hundred of GW bursts would be produced at a $>1$ ns level in the PTA frequency domain. Most of the signals will be washed out in the confusion noise due to the emission of `ordinary' low eccentric binaries. However, their peculiar features may guide the development of targeted data analysis techniques, that may help to recognize them even if overwhelmed by the confusion noise. Employing a minimal criterion for source resolvability (which provides a strict lower limit), we found that less than one system may be actually pinned down at ns precisions. By running several dozens of Monte Carlo realization of the signal from the cosmological population of SMBH binaries and triplets we quantified a statistical $2-50$\\% chance of having a resolvable burst in the Universe (assuming 10 yrs of observation). The probability for detection with {\\rm {\\it LISA}} is essentially nil. However, we stress the fact that we focused on systems with ${\\cal M}>10^7\\msun$; our results then simply imply that it is extremely unlikely that a system which would normally emit outside the {\\it LISA} range will produce a burst in the {\\it LISA} window because of resonant three body interactions. On the other hand, if a consistent fraction of {light} binaries (${\\cal M}<10^7\\msun$) is involved in triple systems, we may expect several eccentricity-driven coalescences to be observed by {\\it LISA}. This eventuality would call for the development of extremely eccentric templates ($e>0.9$) for merging SMBHs, and of adequate analysis techniques to extract the signal."
    },
    "0910/0910.4194_arXiv.txt": {
        "abstract": "We present \\jb, \\hb\\ and \\kb\\ photometry of the Orion Nebula Cluster obtained at the CTIO/Blanco 4~m telescope in Cerro Tololo with the ISPI imager. From the observations we have  assembled a catalog of  about $\\sim7800$ sources distributed over an area of approximately $30\\arcmin\\times40\\arcmin$,   the largest of any survey deeper than 2MASS in this region. The catalog provides absolute coordinates accurate to about 0.15 arcseconds and $3\\sigma$ photometry in the 2MASS  system down to $\\jb\\simeq 19.5$~mag,  $\\hb\\simeq18.0$~mag, $\\kb \\simeq18.5$~mag, enough to detect planetary size objects 1~Myr old under $A_V\\simeq10$~mag of extinction at the distance of the Orion Nebula. We present a preliminary analysis of the catalog, done comparing the (\\jb-\\hb,\\hb-\\kb) color-color diagram,  the (\\hb,\\jb-\\hb) and (\\kb,\\hb-\\kb) color-magnitude diagrams and the \\jb\\hb\\kb\\ luminosity functions of three  regions at increasing projected distance from the Trapezium. Sources in the inner region typically show IR\\ colors compatible with reddened T\\ Tauri stars, whereas the outer fields are dominated by field stars seen through an amount of extinction  which decreases with the  distance from the  center. The color-magnitude diagrams make it possible to clearly distinguish between the main ONC population, spread across the full field,  and background sources. % The luminosity functions of the inner region, corrected for completeness, remain relatively flat  in the sub-stellar regime regardless of the strategy adopted to remove background contamination. ",
        "introduction": "\\label{sec:intro} The Orion Nebula hosts the richest cluster of\\  young ($\\tau\\simeq 1$~Myr) Pre-Main-Sequence (PMS) stars within 1~kpc of the Sun and therefore represents an ideal laboratory to understand the process of star formation \\citep[see][for recent reviews]{Muench+08, ODell+08}. The cluster, both rich ($n\\approx 2000$ members) and dense (about  $2\\times10^4$\\ sources per cubic parsec at its center), is dominated by a small number of massive OB\\ stars mostly clustered in the Trapezium  multiplet ($\\theta^1$~Ori). Their UV\\ emission has created a blister HII\\  \\ region ({\\sl Orion Nebula, M~42, NGC~1976}) whose ionization front is still carving the underlying Orion Molecular Cloud \\cite[OMC-1,][] {ODell+08,ODell+09}. About half of the young cluster members, surrounded by their original circumstellar disks, have been already exposed to the  hard-UV\\ radiation of the Trapezium stars, whereas the other half remain enshrouded within the OMC-1, together with newer active sites of  star formation.  Visible and near-IR\\ data, both in imaging and spectroscopy, are needed to characterize the main physical parameters of the cluster population \\citep{LAH97}. These observations, however, are hampered by the brightness and non-uniformity of the nebular background. To overcome these difficulties, we have exploited the unique combination of sensitivity and spatial resolution offered by  the Hubble Space Telescope (HST)\\ to obtain accurate photometry of the cluster, especially at substellar masses (HST\\ Treasury program GO-10246). Our extensive HST\\ survey has been complemented by ground-based observations, imaging the Orion Nebula from the {\\sl U}-band to the  \\kb-band at La Silla and Cerro Tololo (CTIO)\\ observatories. The ground-based observations, carried out in parallel on the same nights (but at a different epoch than the HST observations), complement the deep HST data which saturate at relatively low brightness levels. Their simultaneity makes the derived stellar colors largely immune to the uncertainties associated with source variability.    The ground-based  survey at visible wavelengths, performed with the Wide Field Camera (WFI) at the ESO/MPG 2.2~m telescope at La Silla, has been recently presented in \\cite{DaRio+09}. In this paper we present the ground-based IR\\ survey, performed at the CTIO/Blanco telescope with the ISPI\\ imager. Its combination of sensitivity and field coverage (about $30\\arcmin\\times40\\arcmin$) ideally complements the previous surveys of this region. In particular, it reaches fainter magnitudes than the wide field survey ($30\\arcmin\\times45\\arcmin$) by Hillenbrand et al. (1998), while it covers a much larger area than the deep surveys done with Keck (\\hb\\kb-bands, $5\\farcm1 \\times 5\\farcm1$, Hillenbrand \\&\\ Carpenter, 2000),  HST-NICMOS\\ (\\jb\\hb-bands, $2\\farcm3\\times2\\farcm3$, Luhman et al. 2000),  UKIRT\\ ({\\sl I}\\jb\\hb-bands,\\ 36 arcmin$^2$, Lucas \\&\\ Roche 2000),  NTT\\ (\\jb\\hb\\kb, $5\\arcmin\\times5\\arcmin$,   Muench et al. 2002) and  Gemini (\\jb\\hb\\kb-bands, 26~arcmin$^2$, Lucas et al. 2005). All these deep surveys concentrate on a field centered around  the Trapezium stars, or in its immediate surroundings in the case of \\cite{LR00}.  In Section\\ \\ref{sec:obs} we discuss the  observing strategy, while in Section\\ \\ref{sec:reduction} we describe the data reduction and photometric calibration. In Section\\ 4 we present our completeness analysis. The resulting source catalog is presented in Section 5, whereas in Section 6 we discuss the color-color and color-magnitude diagrams and the luminosity function derived from our \\jb, \\hb, and \\kb-band photometry. ",
        "conclusions": "We have presented a photometric survey of the Orion Nebula Cluster in the \\jb-, \\hb- and \\kb-passbands carried out at the 4m telescope of Cerro Tololo. The survey, covering a field of about $30\\arcmin\\times40\\arcmin$ centered about $1\\arcmin$\\ southwest of the Trapezium, has been performed in parallel to visible observations of the same region made in La Silla (Da Rio et al. 2009). The two datasets constitute the first panchromatic survey covering simultaneously the Orion Nebula Cluster from the {\\sl U}-band to the \\kb-band. \\   The final catalog, photometrically and astrometrically calibrated to the 2MASS\\  system, contains  7759 sources, (including 174 and 22 sources whose photometry has been taken directly from the 2MASS\\ and \\cite{Mue02} catalogs, respectively). This represents the largest near-IR\\ catalog of the Orion  Nebula Cluster to date. Our sensitivity limits allow to detect  objects of a few Jupiter masses under about $A_V\\simeq 10$ magnitudes of extinction.   We present the color-color diagrams, color-magnitude diagrams and the  luminosity functions for three regions centered on the Trapezium containing the same number of sources (excluding the M43 sub-cluster). Sources in the inner region typically show IR\\ colors compatible with reddened T\\ Tauri stars, whereas the outer fields are dominated by field stars seen through an amount of extinction  which decreases with the projected distance from the  center. The color-magnitude diagrams allow to clearly distinguish between the main ONC population, spread across the full field,  and background sources. After correction for completeness, the luminosity functions in the inner region remains nearly flat, with marginal contamination from background sources."
    },
    "0910/0910.1913_arXiv.txt": {
        "abstract": "We present optical photometry and spectra for the Type Ia supernova (SN Ia) 2007gi in the nearby galaxy NGC 4036. SN 2007gi is characterized by extremely high-velocity (HV) features of the intermediate-mass elements (Si, Ca, and S), with expansion velocities ($v_{\\rm exp}$) approaching $\\sim$15,500 km s$^{-1}$ near maximum brightness (compared to $\\sim$10,600 km s$^{-1}$ for SNe Ia with normal $v_{\\rm exp}$). SN 2007gi reached a $B$-band peak magnitude of 13.25$\\pm$0.04 mag with a decline rate of $\\Delta m_{15}(B)$(true) = 1.33$\\pm$0.09 mag. The $B$-band light curve of SN 2007gi demonstrated an interesting two-stage evolution during the nebular phase, with a decay rate of 1.16$\\pm$0.05 mag (100 days)$^{-1}$ during $t = 60$--90 days and 1.61$\\pm0.04$ mag (100 days)$^{-1}$ thereafter. Such a behavior was also observed in the HV SN Ia 2006X, and might be caused by the interaction between supernova ejecta and circumstellar material (CSM) around HV SNe Ia. Based on a sample of a dozen well-observed $R$-band (or unfiltered) light curves of SNe Ia, we confirm that the HV events may have a faster rise time to maximum than the ones with normal $v_{\\rm exp}$. ",
        "introduction": "Type Ia supernovae (SNe Ia) have been successfully utilized over the past decade to measure the cosmic expansion history \\citep{riess98, per99} and explore the nature of dark energy \\citep[e.g.,][]{riess04, riess07, ast06, wv07}. The foundation for the utility of SNe Ia as a cosmological tool is that some distance-independent observables, such as the light-curve shape parameters \\citep{phi93, riess95, jha07, per97}, the color parameters \\citep{wlf03, wxf05}, or both \\citep{tripp98, guy05}, have been found to correlate with their peak luminosity. These empirical correlations can be used to calibrate the luminosities of SNe Ia and measure their distances with a precision of $\\sim$9\\%.  A recent result suggests that the luminosity standardization of SNe Ia can be improved to a level of $\\sim$6\\% by separating the SNe into two groups based on a spectroscopic criterion [for details, see \\citet{wxf09a}, hereafter W09]. The expansion velocity ($v_{\\rm exp}$) of the SN ejecta is inferred from the blueshift of the absorption minimum of Si`II $\\lambda$6355, and the SNe are divided into one group with normal $v_{\\rm exp}$ (hereafter ``Normal\" SNe Ia) and the other group with high $v_{\\rm exp}$  (hereafter ``HV\" SNe Ia)\\footnote{See \\citet{ben05} and \\citet{bran06} for a similar classification based on the velocity gradient and the strength of the absorption features, respectively.}. W09 found that the two groups have either different extinction laws or color evolution. The cause for such a dichotomy might be related to the properties of their progenitors. The HV SNe Ia are characterized by stronger absorption features of intermediate-mass elements (IMEs, such as Si, S, and Ca) at higher velocities as well as a red $B - V$ color around maximum brightness.  SN 2002bo and SN 2006X are two of the best-studied examples of this class \\citep{ben04, wxf08a}. In particular, the interstellar Na I~D lines were found to show significant variations in the spectra of SN 2006X, likely pointing to the presence of CSM produced by the progenitor system \\citep{pat07, chu08}. The possible detection of CSM around SN 2006X is also supported by a flat evolution of the late-time light curve \\citep{wxf08a} and a detection of light echoes \\citep{wxf08b, cy08}. Similar variability of the Na I~D lines was also observed in SNe 1999cl and 2007le \\citep{blon09, sim09}, two other members of the HV SN Ia class. Contrasting with this, multi-epoch, high-resolution spectral observations of SN 2007af, a Normal SN Ia, do not reveal any significant signature of CSM absorption \\citep{sim07}. This raises the possibility that CSM might be preferentially present for the SNe Ia in the HV class.  In this paper, we present optical observations of another member of the HV SN Ia class, SN 2007gi. Our goal is to increase the sample and understand the properties of well-observed HV SNe Ia. Observation and data reductions are described in \\S 2, while \\S 3 presents the $BVRI$ light curves, color curves, reddening estimate, and an analysis of the rise time. Section~4 presents the spectral evolution. Our discussions and conclusions are given in \\S 5. ",
        "conclusions": "In this paper, we present optical photometry and spectroscopy of SN 2007gi, a SN~Ia with an extremely high expansion velocity measured from IMEs (Si, Ca, and S). We also conduct a comparison study for a sample of SNe~Ia with high and normal expansion velocities (HV and Normal, respectively) to investigate differences in their photometric and spectroscopic behaviors.  The $B$-band light curve of SN 2007gi shows a two-stage evolution: a decay rate of $\\beta = 1.16 \\pm 0.05$ mag (100~days)$^{-1}$ during $t \\approx 60$--90 days and $\\beta = 1.61 \\pm 0.05$ (100~days)$^{-1}$ thereafter. The latter decay rate is similar to those observed in Normal SNe Ia. This two-stage evolution is also present in the HV SN Ia 2006X. More late-time observations of the HV objects are necessary to establish whether this is a universal property. The $B - V$ color of SN 2007gi is found to evolve at a faster slope than that of Normal SNe Ia during the nebular phase, a trend that is also observed for the HV SN Ia 2006X \\citep{wxf08a} and most of the other HV events (Wang, X., et al. 2009c, in preparation). SN 2007gi was detected at $t \\approx 600$ days after $B$ maximum in {\\it HST}/WFPC2 images.  The magnitudes and colors at this very late-time epoch suggest that SN 2007gi may not be contaminated by a light echo on interstellar scales.  Using a dozen well-observed SNe Ia, we confirm the previous claim that HV SNe Ia tend to have a faster rise time than Normal SNe Ia \\citep{pig08}: at the same value of $\\Delta m_{15}$, the rise time of the HV objects is shorter than that of the Normal objects by 1--2 days.  The spectral evolution of SN 2007gi is characterized by high expansion velocities measured from the absorption lines. The value of $v_{\\rm exp}$ is measured to be 15,500 km s$^{-1}$ at $t = -1$ day, $\\sim$50\\% higher than what is observed for Normal SNe Ia. A small flux ratio $\\mathcal{R}$(Si~II) = 0.07$\\pm$0.03 is observed, which traditionally means a hot photospheric temperature in the SN ejecta. Our synthetic spectral analysis using the SYNOW code, on the other hand, suggests that SN 2007gi has a typical photospheric temperature compared with other Normal SNe Ia.  The two-stage evolution in the $B$-band light curve of SN 2007gi may suggest that an additional energy source besides radioactive decay plays a role in the nebular phase.  Dust scattering (a light echo) seems unlikely because the tail luminosity would have remained nearly constant for a long time, which is not the case for SN 2007gi.  Nevertheless, the possibility of a local light echo on a small scale (due to CSM dust) cannot be completely ruled out. Another potential energy source is interaction with CSM; detection of variable sodium lines in the spectra of some HV events provides some evidence for this scenario. Assuming that the CSM lies at a distance of $\\sim10^{17}$ cm from the supernova, it is expected that the outermost ejecta will begin to interact with the CSM at $\\sim$30 days after the $B$ maximum. This naturally accounts for a flatter evolution seen in the $B$-band light curve of the HV SNe Ia after $t = 40$ days. The difficulty with this scenario is the lack of convincing observational evidence for stripped material, such as low-velocity H$\\alpha$ emission in the nebular spectra [e.g., \\citet{mat05}, but see \\citet{leo08} for alternative explanations regarding the observed lack of hydrogen], and the lack of spectral evidence in support of ongoing CSM interaction.  A possible explanation for the HV features observed in the HV SNe Ia is that there are density enhancements of the IMEs in the outer layers of the ejecta of the HV objects (due to a metallicity effect, delayed detonations, or CSM interaction); the $\\gamma$-ray heating is less effective, and the photosphere is formed at an outer layer with a higher expansion velocity compared to the Normal objects."
    },
    "0910/0910.0915_arXiv.txt": {
        "abstract": "We present a pre-discovery $H$-band image of the \\object{HR 8799} planetary system that reveals all three planets in August 2007.  The data were obtained with the Keck adaptive optics system, using angular differential imaging and a coronagraph.  We confirm the physical association of all three planets, including HR~8799d, which had only been detected in 2008 images taken two months apart, and whose association with HR~8799 was least secure until now.  We confirm that the planets are 2--3~mag fainter than field brown dwarfs of comparable near-infrared colors. We note that similar under-luminosity is characteristic of young substellar objects at the L/T spectral type transition, and is likely due to enhanced dust content and non-equilibrium CO/CH$_4$ chemistry in their atmospheres. Finally, we place an upper limit of $\\gtrsim$18~mag per square arc second on the $>$120~AU $H$-band dust-scattered light from the HR~8799 debris disk.  The upper limit on the integrated scattered light flux is $10^{-4}$ times the photospheric level, 24 times fainter than the debris ring around HR~4796A. ",
        "introduction": "The three-planet system around the A5V star HR~8799 \\citep[][henceforth, M08]{marois_etal08b} is the first directly imaged extrasolar multi-planet system.  Along with a reported extrasolar planet around Fomalhaut \\citep[A4V;][]{kalas_etal08}, HR~8799b, c, and d are also the first bona-fide extrasolar planets directly imaged around stars other than the Sun. The HR~8799 and Fomalhaut planetary systems share two important similarities.  First, in both cases the planets are widely separated from their host stars, at projected separations between 24--119~AU.  And second, both hosts are A stars surrounded by cold debris disks \\citep{aumann85, sadakane_nishida86}, circumscribing the planetary systems.  The debris disk around HR~8799 is one of the most massive detected by {\\it IRAS} and is a factor of several brighter \\citep[$L_{\\rm IR}/L_\\ast=2.3\\times10^{-4}$;][]{moor_etal06, rhee_etal07} than that around Fomalhaut \\citep[$L_{\\rm IR}/L_\\ast=8\\times10^{-5}$;][and references therein]{rhee_etal07}.  While Fomalhaut's disk has been spatially resolved at wavelengths ranging from the visible \\citep{kalas_etal05} to the submillimeter \\citep{holland_etal98, marsh_etal05}, the HR~8799 debris disk remained unresolved until recently \\citep[][henceforth, S09]{su_etal09}.  Prior to the discovery of the planets around HR~8799 by M08, we targeted the star with the Keck adaptive optics (AO) system \\citep{wizinowich_etal00} to detect the debris disk in scattered light.  The strength of its IR excess and a detection in the submillimeter \\citep[$\\approx$10~mJy at 850~$\\micron$;][]{williams_andrews06} indicated a substantial optical depth, which given suitable disk viewing geometry could result in a detection of dust-scattered light.  The data presented here did not produce the sought-after scattered light detection, but have revealed all three known planets around HR~8799.  In particular, we report the earliest-epoch image of the closest-in planet, HR~8799d. Prior epoch detections of HR~8799b and/or HR~8799c have been published in M08, \\citet{fukagawa_etal09}, and \\citet{lafreniere_etal09}. ",
        "conclusions": "We have presented an year-2007 image of all three known HR~8799 planets, including a detection of the inner-most planet HR~8799d that precedes its discovery in \\citet{marois_etal08b} by one year. Our data exclude the presence of additional $>$3~\\Mjup\\ outer planets between 68 and 160~AU from the star. We do not detect scattered light from the debris disk, and place an upper limit of $\\sim$18~mag~arcsec$^{-2}$ on its $H$-band surface brightness at 3--6$\\arcsec$ (120--240~AU) from the star. Based on a more robust $H$-band photometric calibration than was possible in \\citet{marois_etal08b}, we % confirm that HR~8799b, c, and d are very under-luminous or redder compared to field brown dwarfs of similar colors or intrinsic luminosities.  We note that this is a distinctive feature of all young substellar objects near the L/T transition, and speculate that it may be due to enhanced dust content and a significant departure from CO/CH$_4$ chemical equilibrium in cool low-surface gravity substellar atmospheres."
    },
    "0910/0910.0853_arXiv.txt": {
        "abstract": "We present the discovery of spectacular double X-ray tails associated with \\ga\\ and a possibly heated X-ray tail associated with \\gab, both late-type galaxies in the closest rich cluster Abell~3627. A deep \\chandra\\ observation of \\ga\\ allows us for the first time to examine the spatial and spectral properties of such X-ray tails in detail. Besides the known bright tail that extends to $\\sim$ 80 kpc from \\ga, a fainter and narrower secondary tail with a similar length was surprisingly revealed, as well as some intriguing substructures in the main tail. There is little temperature variation along both tails. The widths of the secondary tail and the greater part of the main tail also remain nearly constant with the distance from the galaxy. All these results challenge the current simulations. The \\chandra\\ data also reveal 19 X-ray point sources around the X-ray tails. We identified six X-ray point sources as candidates of intracluster ULXs with $L_{\\rm 0.3 - 10 keV}$ of up to 2.5$\\times10^{40}$ erg s$^{-1}$. \\gem\\ spectra of intracluster HII regions downstream of \\ga\\ are also presented, as well as the velocity map of these HII regions that shows the imprint of \\ga's disk rotation. For the first time, we unambiguously know that active star formation can happen in the cold ISM stripped by ICM ram pressure and it may contribute a significant amount of the intracluster light. We also report the discovery of a 40 kpc X-ray tail of another late-type galaxy in A3627, \\gab. Its X-ray tail seems hot, $\\sim$ 2 keV (compared to $\\sim$ 0.8 keV for \\ga's tails). The H$\\alpha$ data for \\gab\\ are also presented. We conclude that the high pressure environment around these two galaxies is important for their bright X-ray tails and the intracluster star formation. The soft X-ray tails can reveal a great deal of the thermal history of the stripped cold ISM in mixing with the hot ICM, which is discussed along with intracluster star formation. ",
        "introduction": "The intracluster medium (ICM) has long been proposed to play a vital role in galaxy evolution in clusters, through ram pressure and turbulent/viscous stripping of the galactic cold gas (e.g., Gunn \\& Gott 1972; Nulsen 1982; Quilis et al. 2000). As the halo gas and the cold interstellar medium (ISM) is depleted in the stripping process, the galactic star formation will eventually be shut down and blue disk galaxies may turn into red galaxies (e.g., Quilis et al. 2000). The removal of the cold ISM also affects the accretion history of the central SMBH. Stripping of the ISM of the disk galaxies in clusters has been extensively studied in simulations recently, with better resolution and more physics included (e.g., Abadi et al. 1999; Stevens et al. 1999; Quilis et al. 2000; Schulz \\& Struck 2001; Bekki \\& Couch 2003; Roediger \\& Hensler 2005; Kapferer et al. 2009). These simulations show that stripping has a significant impact on galaxy evolution (e.g., disk truncation, formation of flocculent arms, build-up of a central bulge and enhanced star formation in the inner disk at the early stage of the interaction).  Besides the impact on galaxy evolution, another significant question related to stripping is the evolution of the stripped ISM. After the cold ISM is removed from the galaxy, the general wisdom is that the stripped gas is mixed with the ICM through evaporation eventually. However, it is now known that a fraction of the stripped ISM turns into new stars in the galactic halo or the intracluster space, as revealed from observations (e.g., Sun et al. 2007b, S07 hereafter) and simulations (e.g., Vollmer et al. 2001b; Schulz \\& Struck 2001; Kronberger et al. 2008; Kapferer et al. 2009). The stripped cold ISM, if it can survive long enough to reach the cluster center, can effectively heat the cluster core via ram-pressure drag, which makes gravitational heating by accretion a possible way to offset cooling (Dekel \\& Birnboim 2008 argued for 10$^{5}$ - 10$^{8}$ M$_{\\odot}$ clumps). Current HI observations of cluster spiral galaxies still fail to detect most of the HI gas missing from HI-deficient spiral galaxies in the intracluster space (e.g., Vollmer \\& Huchtmeier 2007), which implies that the bulk of the stripped HI gas has been heated out of the cold phase. However, little is known about the details of mixing, as the actual strength of heat conductivity and viscosity is poorly known. How quickly is the stripped cold ISM heated and mixed with the ICM? What observational signature will evaporation and mixing produce? What fraction of the stripped cold ISM turns into new stars? The mixing of the stripped cold ISM with the hot ICM will produce multi-phase gas. Depending on the poorly known details of mixing, prominent soft X-ray emission may be produced, as well as H$\\alpha$ emission. In this ``unified model'' for tails of cluster late-type galaxies, X-ray and H$\\alpha$ tails are simply manifestations of the cold ISM tail in the mixing process. Therefore, one can better understand mixing and the relevant micro-physics from multi-wavelength data, e.g., the amount of missing HI gas and the amount of the stripped gas in hotter phases. Obviously, a central problem is the energy transfer in the multi-phase gas, which is also a significant question for large cool cores in clusters. Moreover, the wake behind the galaxy produced by stripping provides a way to constrain the ICM viscosity (e.g., Sun et al. 2006, S06 hereafter; Roediger \\& Br$\\ddot{\\rm u}$ggen 2008a).  Various dynamic features, e.g., bow shocks, tails and vortices, can be produced in stripping (e.g., Stevens et al. 1999; Schulz \\& Struck 2001; Roediger et al. 2006; Roediger \\& Br$\\ddot{\\rm u}$ggen 2008b, RB08 hereafter; Tonnesen \\& Bryan 2009). Observational evidence of stripping of cluster late-type galaxies is present in HI and H$\\alpha$ observations, either through HI deficiency or tails (e.g., Giovanelli \\& Haynes 1985; Gavazzi et al. 2001b; Kenney et al. 2004; Oosterloo \\& van Gorkom 2005; Chung et al. 2007; Yagi et al. 2007; S07). X-ray tails of late-type cluster galaxies have only been detected recently: C153 in A2125 at $z$=0.253 (Wang et al. 2004), and UGC~6697 in A1367 at $z$=0.022 (Sun \\& Vikhlinin 2005). However, there are only $\\sim$ 60 counts from C153 so the X-ray extension is not unambiguous (Wang et al. 2004). The \\xmm\\ data show that the X-ray tail of UGC~6697 may extend to $\\sim$ 90 kpc from the nucleus but most of the X-ray emission is within the galaxy. Moreover, the galaxy has a peculiar morphology and a tidal tail. Its complex kinematic behavior, revealed from the velocity map, implies the existence of a second galaxy hidden behind the main galaxy (Gavazzi et al. 2001a). Thus, UGC~6697 is a very complicated system where tidal interaction is important. There is also a wide tail (67 kpc$\\times$90 kpc) of the spiral NGC~6872 in the 0.5 keV Pavo group (Machacek et al. 2005), but it is uncertain how it was formed as the tail terminates on the dominant early-type galaxy of the group. The brightest known X-ray tail behind a cluster late-type galaxy is \\ga\\ in the closest rich cluster A3627 (S06), with $\\sim$ 80\\% of the X-ray emission beyond the galactic halo. One should be aware that strong X-ray tails of late-type galaxies are rare (e.g., none in Virgo, see Section 7.1). We have done a blind search of strong X-ray tails in nearby clusters ($z<0.06$) by extending our previous \\chandra\\ work (Sun et al. 2007a) to the current \\xmm\\ and \\chandra\\ archives. Despite extensive cluster data in the archives, e.g., mosaic fields of $\\sim$ 2 deg$^{2}$ around the Perseus and Coma clusters with \\xmm, only two more X-ray tails of cluster late-type galaxies were found. One is also in A3627 (\\gab) and will be discussed in this paper. The other one is NGC~4848 in the Coma cluster (Finoguenov et al. 2004) that was recently observed for 29 ks by \\chandra. However, both X-ray tails are shorter (40 - 50 kpc) and 2.4 - 5.5 times fainter than \\ga's tail. Thus, the proximity of \\ga, the high flux and the large length of its X-ray tail makes it the best target for detailed analysis and comparison with simulations. One should not confuse X-ray tails of early-type galaxies with those of late-type galaxies. Early-type galaxies in clusters have abundant X-ray ISM but little or no cold ISM. There are some X-ray tails of early-type galaxies reported (e.g., Machacek et al. 2006; Sun et al. 2007; Randall et al. 2009). Their X-ray tails are composed of the stripped hot ISM and do not co-exist with cold gas.  A3627 is the closest massive cluster ($z$=0.0163, $\\sigma_{\\rm radial}$ = 925 km/s and $kT$=6 keV) rivaling Coma and Perseus in mass and galaxy content (Kraan-Korteweg et al. 1996; Woudt et al. 2008). It is also the sixth brightest cluster in RASS and the second brightest non-cool-core cluster after Coma (B$\\ddot{\\rm o}$hringer et al. 1996). \\ga\\ is a blue emission-line galaxy (Woudt et al. 2004) that is only $\\sim$ 180 kpc from the cluster's X-ray peak in projection. Its radial velocity (4680$\\pm$71 km~s$^{-1}$, Woudt et al. 2004) is close to the average velocity in A3627 (4871$\\pm$54 km~s$^{-1}$, Woudt et al. 2008) so most of its motion is probably in the plane of sky. If \\ga\\ is on a radial orbit, its real distance to the cluster center should be close to the projected distance. There is no sign of a galaxy merger and no strong tidal features around \\ga. S06 found a long X-ray tail behind \\ga, in both \\chandra\\ (14 ks) and \\xmm\\ data (MOS: 18 ks, PN: 12 ks). The tail extends to at least 70 kpc from the galaxy with a length-to-width ratio of $\\sim$ 10. The X-ray tail is luminous ($\\sim$ 10$^{41}$ erg s$^{-1}$), with an X-ray gas mass of $\\sim 10^{9}$ M$_{\\odot}$. S06 interpreted the tail as the stripped ISM of \\ga\\ mixed with the hot ICM. The \\chandra\\ data also reveal three hard X-ray point sources ($L_{X} \\sim 10^{40}$ erg s$^{-1}$) along the tail, and the possibility of all of them being background active galactic nuclei (AGNs) is very small. S06 suggested that some of them may be ultra-luminous X-ray sources (ULXs) born from active star formation in the tail. S07 further discovered a 40 kpc H$\\alpha$ tail and over 30 emission-line regions downstream of \\ga, up to 40 kpc from the galaxy. S07 concluded that they are giant HII regions in the halo downstream or in intracluster space. Sivanandam et al. (2010) observed \\ga\\ with IRAC and IRS on \\spi. A warm ($\\sim$ 160 K) molecular hydrogen tail was detected to at least 20 kpc from the galaxy from the IRS data, at the same position as the bright X-ray and H$\\alpha$ tail. The total mass of the warm H$_{2}$ gas is $\\sim 2.5\\times10^{7}$ M$_{\\odot}$. As the IRS fields only cover the front part of the X-ray tail, the above warm H$_{2}$ gas mass is only a lower limit. 8 $\\mu$m PAH emission is also detected from the bright HII regions identified by S07. The \\spi\\ results may imply a large reservoir of colder gas than the observed warm gas downstream of \\ga, coexisting with the hot ICM ($\\sim 7\\times10^{7}$ K) and the H$\\alpha$ emitting gas ($\\sim 10^{4}$ K).  Obvious follow-ups include deep \\chandra\\ exposures and optical spectroscopic observations of the candidate HII regions. In this paper, we present the results from our deep \\chandra\\ observation of \\ga\\ and \\gem\\ GMOS spectroscopic observations. The plan of this paper is as follows: The \\chandra\\ and \\gem\\ observations and the data analyses are presented in Section 2. In Section 3, we discuss the X-ray tails of \\ga. Section 4 is on the properties of the surrounding ICM. Section 5 discusses the \\chandra\\ point sources and the intracluster HII regions. We also found another 40 kpc X-ray tail in A3627, associated with \\gab, which is discussed in Section 6. The results are discussed in Section 7 and Section 8 contains the summary. We adopt a cluster redshift of 0.0163 for A3627 (Woudt et al. 2008). Assuming H$_{0}$ = 71 km s$^{-1}$ Mpc$^{-1}$, $\\Omega$$_{\\rm M}$=0.27, and $\\Omega_{\\rm \\Lambda}$=0.73, the luminosity distance is 69.6 Mpc, and 1$''$=0.327 kpc. ",
        "conclusions": "\\subsection{X-ray Tails of Late-type Galaxies and Mixing}  While the paucity of luminous X-ray tails of late-type galaxies in nearby clusters has been known (e.g., Sun et al. 2007a), the discovery of \\ga's double X-ray tails is a surprise. M86 also has double X-ray tails and Randall et al. (2008) interpreted it as the consequence of the aspherical potential of M86 and M86's orbit. M86 is an early-type galaxy so its potential structure is different from \\ga's and its tail is only composed of the stripped hot ISM. M86's tails are in fact connected to the ``plume'' on the north of the galaxy and are located at larger distance from the galaxy, instead of coming from the galaxy directly like \\ga's. Moreover, a spectacular H$\\alpha$ complex connecting M86 and the nearby disturbed spiral NGC~4438 was recently discovered by Kenney et al. (2008). Some H$\\alpha$ emission also extends to the position of M86's tail. Thus, tidal interaction may also be important for M86's X-ray tail. There also appears to be a split feature that resembles a double tail in NGC~4438's \\chandra\\ image (Machacek et al. 2004; Randall et al. 2008). However, one part of the split feature is almost along the distorted disk plane and tidal interaction should be important in this case (Machacek et al. 2004; Kenney et al. 2008). NGC~4438 is $\\sim$ 3 times more luminous than \\ga\\ in the $K_{\\rm s}$ band. Its truncated halo should be larger than \\ga's ($\\sim$ 15 kpc in radius, S07), also because the Virgo cluster is less massive than A3627. The features in NGC~4438 only extend to $\\sim$ 11 kpc from the nucleus, clearly within its halo. The other late-type galaxy in the proximity, NGC~4388, has many H$\\alpha$ filaments and a spectacular HI tail (Yoshida et al. 2002; Oosterloo \\& van Gorkom 2005). A faint X-ray extension can be traced to $\\sim$ 14 kpc from the nucleus in the direction of the HI tail, from the \\chandra\\ and \\xmm\\ data. There is also an X-ray extension to east in the disk plane, which creates the appearance of double tails (Randall et al. 2008). NGC~4388 is $\\sim$ 1.5 times more luminous than \\ga\\ in the $K_{\\rm s}$ band so similar to NGC~4438, its X-ray tail is still within its halo. Thus, we do not consider the one-sided X-ray features in NGC~4438 and NGC~4388 to be X-ray double tails based upon the current data. They are clearly much fainter than \\ga's tails, especially given \\ga's much larger distance ($\\sim$ 4.3 times) and the much higher absorption to the direction of \\ga. We can also define strong X-ray tails as one-sided features outside of the tidal truncation radius of the galactic halo (like the tails of \\ga\\ and \\gab), in order to distinguish them from the weak X-ray tails of NGC~4438 and NGC~4388.  The observed X-ray emission of the tails is likely from the interfaces of the hot ICM and the cold stripped ISM, reflecting the multi-phase nature of mixing. Although the detailed mixing process that determines the emergent X-ray spectra is poorly understood, the observed tail temperature may be related to the mass-weighted temperature that is usually adopted in simulations, $T_{\\rm mw} = (M_{\\rm ISM} T_{\\rm ISM} + M_{\\rm ICM} T_{\\rm ICM}) / (M_{\\rm ISM} + M_{\\rm ICM}) \\approx T_{\\rm ICM} / (1+X)$, where $X = M_{\\rm ISM} / M_{\\rm ICM}$ and we assume that $T_{\\rm ISM} \\ll T_{\\rm ICM}$. For the X-ray tails of \\ga\\ and \\gab, $X = 2 - 6.5$ if we assume $T_{\\rm mw}$ is the observed spectroscopic temperature. However, it is easy to find out that the mass-weighted temperature is always larger than the spectroscopic temperature derived by fitting the observed spectra with a single-$kT$ model ($T_{\\rm spe}$). We examined this with the CEMEKL and APEC models. For the CEMEKL model, $d EM(T) \\propto (T/T_{\\rm max})^{\\alpha-1} d T$. Assuming an isobaric condition and for $\\alpha >$ -1,  \\begin{equation} T_{\\rm mw} = \\frac{\\int T dm(T)}{\\int dm(T)} = \\frac{1+\\alpha}{2+\\alpha} T_{\\rm max} \\end{equation}  For $\\alpha >$ -0.5, 1/3 $T_{\\rm max} < T_{\\rm mw} < T_{\\rm max}$. It is also reasonable to assume that $T_{\\rm max} \\leq T_{\\rm ICM}$. However, with a series of XSPEC simulations, we found that $T_{\\rm spe} < T_{\\rm mw}$. For example, $T_{\\rm spe}$ = 0.5 - 1 keV when -0.5 $< \\alpha <$ 0.5, while $T_{\\rm mw}$ = 2 - 3.6 keV for $T_{\\rm max}$ = 6 keV. This is not surprising as $T_{\\rm spe}$ is mainly determined by the centroid of the iron-L hump at low temperatures, which biases $T_{\\rm spe}$ low. As long as there are significant emission components at $kT$ = 0.4 - 2 keV, the iron-L hump will be strong to make $T_{\\rm spe}$ deviated from $T_{\\rm mw}$. On the other hand, the parameter $\\alpha$ in the CEMEKL model may reflect the age of the X-ray tail. Our simulations indeed show that $T_{\\rm spe}$ increases monotonously with $\\alpha$. From the CEMEKL fits to the spectra of \\ga\\ and \\gab, \\gab's stripping event may indeed be older. The above simple estimate implies that $T_{\\rm spe}$ is a certain fraction of the ICM temperature. The luminosity of the X-ray tail should also increase with the magnitude of the ambient ICM pressure, under an isobaric condition. Thus, the X-ray tails of late-type galaxies in groups and the outskirts of poor clusters may have temperatures (e.g., $<$ 0.4 keV) and densities too low to be bright in the \\chandra\\ and \\xmm\\ energy bands. This helps to explain the paucity of strong X-ray tails of late-type galaxies in nearby clusters (e.g., Sun et al. 2007a). Other factors include the generally high cluster background and the small number of late-type galaxies in the \\chandra\\ field of nearby clusters as the covered cluster volume is small.  With the above scenario, the X-ray tail of a late-type galaxy is simply the manifestation of its cold ISM tail embedded in the hot ICM. Optical line emission is also expected and is indeed detected, e.g., the H$\\alpha$ tails of \\ga\\ and \\gab\\ (S07 and Section 6). HI observations of A3627 were made with Australia Telescope Compact Array (ATCA) in 1996 and no HI emission was detected from \\ga\\ or \\gab\\ (Vollmer et al. 2001a). The 3$\\sigma$ detection limit is $\\sim$ 3 mJy/beam in one velocity channel with a beamsize of 30$''$ (Vollmer et al. 2001a). Adjusting to the cluster distance used in this work and assuming a linewidth of 150 km/s (Vollmer et al. 2001a), the corresponding HI gas mass limit is $\\sim 5\\times10^{8}$ M$_{\\odot}$ per beam. While this limit is not weak for the remaining HI gas in the disks of \\ga\\ and \\gab\\ (roughly within one beam), the X-ray tails of \\ga\\ and \\gab\\ span over 7 - 15 times the beam size so a lot of diffuse HI gas can remain undetected downstream of the galaxy. In fact, with the sensitivity of the available ATCA data (an HI column density limit of 2$\\times10^{20}$ cm$^{-2}$), four of the seven HI tails in the Virgo cluster found by Chung et al. (2007) and the long HI tail of NGC~4388 (Oosterloo \\& van Gorkom 2005) would not be detected. Clearly deeper HI observations are required, although the nearby bright radio source PKS~1610-60 (43 Jy at the 1.4 GHz) makes the task difficult (8$'$ - 14$'$ from \\ga\\ and \\gab). Both \\ga\\ and \\gab\\ are expected to host a significant amount of atomic and molecular gas initially. From Blanton \\& Moustakas (2009), the total neutral plus molecular gas fraction (relative to the stellar mass) is $\\sim$ 70\\% for \\ga\\ and $\\sim$ 30\\% for \\gab. Assuming a type of Sc and a blue-band diameter of 20 kpc, the expected HI mass of \\ga\\ is $\\sim 2.4\\times10^{9}$ M$_{\\odot}$, from the empirical relation derived by Gavazzi et al. (2005). Similarly for \\gab\\ (a type of Sa and a blue-band diameter of 30 kpc), the expected HI mass is $\\sim 3.9\\times10^{9}$ M$_{\\odot}$. Thus, \\ga\\ is expected to have $\\sim (3-5)\\times10^{9}$ M$_{\\odot}$ of cold ISM initially, while \\gab\\ is expected to have $\\sim (4-9)\\times10^{9}$ M$_{\\odot}$ initially. These numbers can be compared with the X-ray gas mass in their X-ray tails (Table 2). Sivanandam et al. (2010) detected the mid-IR emission from warm H$_{2}$ gas ($\\sim$ 150 K) in the galaxy and the first 20 kpc of the main tail, with a total mass of $\\sim 2.5\\times10^{7}$ M$_{\\odot}$. The H$_{2}$ lines that IRS detects are rotationally excited lines that are generated very efficiently. The warm gas they probe is likely a small fraction of the total H$_{2}$ gas, if the bulk of the molecular gas is still cold. Deep HI and CO observations are required to recover the bulk of the cold gas in \\ga, \\gab\\ and their wakes.  The same question can also be asked to other cluster late-type galaxies with signs of stripping. Does a stripped ISM tail show up in HI, H$\\alpha$ and X-rays simultaneously? We attempt to summarize the known examples of stripped tails for cluster late-type galaxies. A Venn diagram is plotted in Figure 15. While a lot of stripped tails are known for cluster late-type galaxies, there is currently little known overlap between different bands, especially between X-ray tails and HI tails. This may be due to the lack of sufficient data, either in X-rays (e.g., for the Virgo galaxies with HI tails, Chung et al. 2007) or in HI (e.g., the two galaxies in this paper). We are pursuing \\xmm\\ data for the Chung et al. (2007) galaxies and deeper ATCA data for the A3627 galaxies to better address this question. On the other hand, there are various reasons that A one-to-one correlation may not exist in every case, even with better data. Depending on its age and the surrounding ICM, an ISM tail may emit predominantly in the HI band or the X-ray band. Soft X-ray emission of the tail may only be bright in high-pressure environment. We also understand little about the details of mixing. It may be related to the mean free path of particles and the coherence length of the magnetic field, which would introduce a radial dependence on the efficiency of mixing. Nevertheless, we first need better data to update the Venn diagram shown in Figure 15.  Deep X-ray data alone, like \\ga's, can reveal some thermal history of the stripped ISM. One surprise from \\ga's X-ray data is the constancy of the spectroscopic temperatures along both tails (Figure 7). Indeed extra care is required to interpret this spectroscopic temperature as it is mainly determined by the centroid of the combined iron-L humps from multi-$T$ gas. Nevertheless, the possible spectral difference between \\ga's tails and \\gab's tail is intriguing. This difference, if confirmed, combined with the temperature constancy of \\ga's tails, puts a constraint on the heating timescale of the stripped cold ISM. Clearly the knowledge of the HI/CO gas distribution in these two galaxies will help to tighten the constraint, although it is certain that this timescale depends on the ICM temperature.  \\subsection{Comparison with Simulations}  After the initial analytic work by Gunn \\& Gott (1972), many simulations have been run on the ram-pressure stripping of late-type galaxies in clusters (e.g., Abadi et al. 1999; Schulz \\& Struck 2001; Vollmer et al. 2001b; RB08; Kapferer et al. 2009; Tonnesen \\& Bryan 2009 and references in those papers). Recent efforts surpass the early ones in many respects, e.g., resolution, time steps, realistic treatment of the galaxy orbit and some important baryon physics (e.g., cooling and star formation). We mainly compare our results with recent simulations of three groups, the Innsbruck group with the GADGET-2/SPH code (Kronberger et al. 2008; Kapferer et al. 2009), the Bremen group with the FLASH/AMR code (Roediger et al. 2006; RB08) and the Columbia group with the {\\em Enzo}/AMR code (Tonnesen \\& Bryan 2009). The GADGET-2 simulations include recipes for cooling, star formation and stellar feedback. The model galaxies run through a wind tunnel with a constant pressure. The FLASH simulations are non-radiative and model the flight of a disk galaxy through a galaxy cluster with realistic ICM and potential distributions. The {\\em Enzo} simulations include cooling but not feedback (Tonnesen \\& Bryan 2009). To mimic effects of some heating processes not in simulations, they impose a minimum temperature floor. The model galaxy is also subject to a constant ram pressure. The inclusion of cooling is important, as it indeed happens as revealed from our data (the intracluster HII regions and star clusters) and it has profound effects in simulations (Kapferer et al. 2009; Tonnesen \\& Bryan 2009). Cooling also naturally adjusts the ISM distribution on the disk plane, creating multi-phase ISM with enhancement and holes (Kapferer et al. 2009; Tonnesen \\& Bryan 2009), which is not present in the non-radiative runs. However, cooling needs to be offsetted by stellar feedback. Unknown efficiencies of transport processes (viscosity and heat conduction) because of unclear strength and configuration of magnetic field, plus the other uncertainties in the prescriptions of star formation and stellar feedback, make the tasks difficult (e.g., Tonnesen \\& Bryan 2009). We emphasize that comparison at the current stage is not straightforward as galaxies in these simulations are always more massive than \\ga\\ and high pressure environment was not explored in all simulations. We also caution that it is best to compare the observational data with the mock data generated from simulations. Kapferer et al. (2009) and Tonnesen \\& Bryan (2009) attempted to create X-ray mock data but more work is required, e.g., folding with the actual \\chandra\\ response, projection on the ICM background and the use of spectroscopic temperature rather than mass-weighted temperature (Section 7.1).  Tails generally appear longer and wider in simulations than in observations. The simulations by Kapferer et al. (2009) show that the X-ray tails can be traced to $\\sim$ 300 kpc from a galaxy that is about 4 times more massive than \\ga, with an ambient pressure similar to \\ga's. The simulated X-ray tails are also very clumpy, with widths generally increasing with the distance from the galaxy. The FLASH simulations also produce longer tails than observed in this work, but are difficult to compare with the X-ray data directly. As RB08 pointed out, the observed tail length depends on the mass-loss per orbital length, which is the highest in poor clusters as galaxies in rich clusters move too fast and the ram pressure in groups is not high enough. The FLASH simulations also show tail flaring (i.e. increasing width with distance from the galaxy downstream) in all cases. The velocity perpendicular to the direction of the ram pressure comes from the random motion of the ICM, the disk rotation and turbulence in the wake. For a rotation velocity of 100 km/s over 50 Myr, the offset is 5 kpc, which is not small in \\ga's case as its tails are narrow. The Tonnesen \\& Bryan (2009) simulations produce long but narrower tails. The tails are also very clumpy. None of these simulations can reproduce the double tails of \\ga. The simulated tails are generally too clumpy with strong turbulence. This simple comparison may point to a higher ICM viscosity than present in the simulations. The Reynold number is $\\sim 3{\\cal M}(L/\\lambda)$, where ${\\cal M}$ is the Mach number, $L$ is the size of the remaining galactic gas disk and $\\lambda$ is the mean free path of particles in the ICM. For unmagnetized gas, $\\lambda \\sim$ 10 kpc around \\ga. The radius of the remaining H$\\alpha$ disk and the galactic X-ray emission is only $\\sim$ 1.7 kpc (S07). Thus, the Reynold number is on the order of unity which implies a laminar flow. Of course the presence of magnetic field would increase the Reynold number. In the simulations by Kapferer et al. (2009), a bright spot in the X-ray images is always found at distances of a few tens of kpc from the galaxy, caused by compressional heating. This feature can be compared with the bright spot in the main tail (Figure 1 and 5). However, the location of this feature is generally too far from the galaxy in simulations with high ambient pressure. Such a bright spot is also not observed in the secondary tail.  The most obvious question about \\ga's X-ray tails is why are there double tails. The simplest explanation is to involve another dwarf galaxy forming a pair with \\ga. This was in fact simulated by Kapferer et al. (2008), showing double gaseous tails, stellar tidal tails and a stellar bridge between two galaxies. The X-ray gas mass of the secondary tail is 40\\% - 50\\% of the mass of the main tail, assuming the same filling factors of the X-ray gas. If we simply take the ratio of their X-ray gas masses as the ratio of their ISM gas masses, the stellar mass of this assumed dwarf that is responsible for the secondary tail should be $\\sim$ 30\\% of \\ga's stellar mass, from the ISM mass - stellar mass relation summarized by Blanton \\& Moustakas (2009). Such a galaxy should be easily detected, at least through the tidal tails and bridge it causes even if it hides behind \\ga. However, the SOAR images do not show the presence of another galaxy near \\ga. There are also no tidal features upstream of the galaxy. There are some continuum features downstream of the galaxy as shown in S07. However, all of them are around confirmed HII regions and have blue colors. Our new \\hst\\ data only reinforce the conclusion from the SOAR images. The central region of the galaxy is very dusty but the galaxy appears regular without significant tidal features in the \\hst\\ $I$ band (Paper IV, in preparation). Thus, we conclude that both X-ray tails originate from \\ga.  The discovery of two separated stripping tails from one galaxy presents a challenge to our understanding of stripping. In simulations (e.g., RB08; Kapferer et al. 2009; Tonnesen \\& Bryan 2009), cold ISM is removed from the sides of the disk or through holes and is quickly mixed to form one clumpy tail because of disk rotation, momentum transferred from the ICM and turbulence in the wake, especially if projected on the plane of sky. Perhaps we should begin to consider more realistic ISM distributions in the galaxy. Optical images clearly show the presence of two thick spiral arms, one to the north and the other one to the south (S07; better shown in the \\hst\\ images). Maybe the two X-ray tails correspond to stripping of molecular gas concentrated around these two spiral arms. CO images of nearby face-on spiral galaxies indeed show that molecular gas distribution follows major spiral arms in many cases (e.g., Kuno et al. 2007). This is also generally true for HI gas (e.g., Adler \\& Westpfahl 1996). For \\ga, we present a simple cartoon in Figure 16. The north arm is apparently the trailing arm where stripping is easier so a brighter X-ray tail is produced. The south arm is the leading arm so gas removal there is more difficult as most gas would be pushed inwards to deeper galactic potential first. This also explains the fact that more intracluster HII regions (at $>$ 15 kpc from the nucleus) are found around the main tail while the numbers of HII regions in the halo are similar at two sides. The curvature of the secondary tail may come from the galactic rotation, which is high enough to cause the observed offset. The main tail may also be curved in 3D but appears straight in projection. The rather regular wake structure may imply significant viscosity. On the other hand, draping of the ICM magnetic field (e.g., Dursi \\& Pfrommer 2008) can also inhibit the motion of the stripped ISM perpendicular to the direction of ram pressure.  In summary, it remains to be seen whether simulations can reproduce the double X-ray tails of \\ga\\ as observed, and the gas properties in other bands (H$\\alpha$ and H$_{2}$ at least). We also emphasize that the comparison between data and simulations is not trivial and straightforward. Observationally we need deeper data at multiple wavelengths. For simulations, mock data need to be produced with sensitivities matching the actual observations. Uncertainties in the viscosity, heat conduction, cooling and stellar feedback should be explored with different runs. Eventually we need to compare the multi-wavelength data (HI, CO, IR, optical and X-rays) with the mock data in the same band. Besides the simple comparison of morphology, a lot of details of the tail properties (e.g., temperature distribution, filling factors of the gas in different phases and energy transfer between gas in different phases) can be better examined.  \\subsection{Formation of Young Stars and X-ray Binaries in the Stripped ISM}  The new observations of \\ga, for the first time, unambiguously confirm the active star formation in the cold ISM stripped by the ICM ram pressure. As emphasized in S07, \\ga's immediate surroundings are devoid of bright galaxies. There are only two galaxies that are within 4 mag of \\ga\\ in the $I$ band, within 100 kpc from the end of \\ga's X-ray tail (G1 and G2 in Figure 1 of S07). Their redshifts are unknown and they are 1.9 - 2.7 mag fainter than \\ga\\ in the $I$ band. There is also no such galaxy within 70 kpc from the nucleus of \\ga. Both G1 and G2 are undisturbed and there are no tidal features between \\ga\\ and G1 (or G2). Our \\hst\\ data do not reveal another dwarf galaxy near \\ga\\ (Paper IV, in preparation). The velocity map of the HII regions also shows no extra component, besides \\ga's rotation pattern. It is also known that the cluster potential is not the main factor for stripping (S07; Sivanandam et al. 2010), although it indeed affects the size of \\ga's dark matter halo. Thus, we conclude that ICM ram pressure is responsible for the stripping of the ISM. In our case, the ambient pressure is even high enough to remove some molecular gas from the galaxy. The HI gas has typical pressures of $P/k_{\\rm B} \\sim 10^{3} - 10^{4}$ K cm$^{-3}$, while the pressure in the core of a starless giant molecular cloud can be higher than $10^{6}$ K cm$^{-3}$. The ambient total thermal pressure (electron + ions) is 1.3$\\times10^{5}$ K cm$^{-3}$ for $n_{\\rm e}$ = 10$^{-3}$ cm$^{-3}$ and $kT$=6 keV, while the ram pressure is 3.2$\\times10^{5}$ K cm$^{-3}$ for $n_{\\rm e}$ = 10$^{-3}$ cm$^{-3}$ and $v_{\\rm gal}$ = 1500 km/s.  After removing the cold ISM from the galaxy, the key for the intracluster star formation is that the stripped ISM must be able to cool, likely aided by shocks and the absence of the UV radiation field that is strong on the disk plane. Intracluster star formation in the stripped ISM was first suggested in simulations (Schulz \\& Struck 2001; Vollmer et al. 2001b). Schulz \\& Struck (2001) even suggested formation of dwarf galaxies in this mode. In the new GADGET-2/SPH simulations by Kronberger et al. (2008) and Kapferer et al. (2009), new stars are formed in the stripped ISM, forming $\\sim$ 1 kpc structures with masses of up to 10$^{7}$ M$_{\\odot}$ individually to hundreds of kpc downstream of the galaxy. However, all simulations suffer from uncertainties in transport processes like viscosity and heat conduction, which are important for determining the heating rate.  More evidence of intracluster star formation, either HII regions or blue star clusters, has been revealed recently (e.g., Lee et al. 2000; Gerhard et al. 2002; Cortese et al. 2004; Cortese et al. 2007; Yoshida et al. 2008). They are only at one side of the galaxy. In some cases, tidal interaction may be important (e.g., Lee et al. 2000). But they all lack unambiguous associations with gaseous tails. We especially highlight the growing number of examples of one-sided blue stellar filaments and star clusters behind cluster galaxies, which includes two galaxies in A2667 and A1689 at $z \\sim$ 0.2 (Cortese et al. 2007) and a dwarf galaxy RB~199 in the Coma cluster (Yoshida et al. 2008). The connection between the first two galaxies at $z \\sim 0.2$ with \\ga\\ was discussed in S07. In the case of RB~199, Yoshida et al. (2008) found narrow blue filaments, knots, H$\\alpha$ filaments and clouds extending to 80 kpc from the galaxy in one side. It is a rather remarkable case as the galaxy is much fainter than all previous galaxies, with a stellar mass of $\\sim 8 \\times 10^{8}$ M$_{\\odot}$. The stellar mass of the downstream structures is $\\sim 10^{8}$ M$_{\\odot}$, which is $\\sim$ 10\\% of the galactic mass. Yoshida et al. (2008) also favored a mechanism similar to what we suggested for \\ga\\ in S07. These three examples may imply a rather common evolution stage for cluster late-type galaxies. As the stage with bright blue stellar trails is short ($\\leq$ 200 Myr), more similar examples are expected only in large surveys.  If we assume that intracluster star formation happens in the stripped ISM of every cluster spiral, can this mode contribute a significant fraction of the intracluster light? The answer depends on the mass fraction of intracluster stars formed out of the stripped ISM (or the star formation efficiency). S07 gave a rough estimate of $\\sim$ 1\\%, which may only be a lower estimate as blue stellar filaments and clusters (not shining in H$\\alpha$) are not counted. While a detailed account of the total stellar mass downstream of \\ga\\ is the aim of Paper IV, here we take 1\\% as a lower limit. If the efficiency of the intracluster star formation is only 1\\%, the intracluster light contributed from star formation in stripped ISM is small, given the mass fraction of the intracluster light to the total stellar light in clusters ($\\sim$ 20\\%, likely depending on the halo mass, e.g., Krick \\& Bernstein 2007). However, the case of RB~199 implies a fraction of $\\sim$ 10\\%, if the original ISM in RB~199 had a mass comparable to its stellar mass (e.g., Blanton \\& Moustakas 2009). In the simulations by Kronberger et al. (2008) and Kapferer et al. (2009), 7\\% - 12\\% of the initial ISM turns into new stars in the wake. If an efficiency of 10\\% is adopted, the intracluster light contributed by ram pressure stripping is a significant fraction of the total intracluster light. However, this efficiency may not be too high given the star formation efficiency in galaxies. In galactic disks, $\\sim$ 1\\% of gas is converted to stars per free-fall time (e.g., Leroy et al. 2008). In 100 Myr (a typical orbital time in the disks), $\\sim$ 6\\% of gas is converted to stars (e.g., Kennicutt 1998; Leroy et al. 2008). We note that 100 Myr is also about the orbital time of the residual H$\\alpha$ disk ($\\sim$ 1.7 kpc radius with an orbital velocity of $\\sim$ 100 km/s) and close to the age of the X-ray tails. Clearly, a lot of work is required to constrain the intracluster star formation efficiency. The efficiency may also change with the mass of the host galaxy as more stars are formed inside the halo of more massive galaxies. In principle, the ratio of the intracluster SNIa to the other types can also inform us the significance of intracluster star formation.  As active star formation is confirmed downstream of the galaxy, the discovery of ULXs is not a surprise. As we estimated in Section 5.2, the total luminosity of ULXs can be consistent with the expectation from the intracluster star formation rate. The space density of X-ray point sources also traces the strength of the star formation, stronger in the halo and weaker around the unbound tail. A similar system may be the high-velocity system of NGC~1275. Gonzalez-Martin et al. (2006) found a concentration of eight ULXs ($L_{\\rm 0.5-7 keV} > 3\\times10^{39}$ erg s$^{-1}$) around the high-velocity system. They may be associated with the active star formation in and downstream of the high-velocity system, similar to what is happening in \\ga."
    },
    "0910/0910.0254_arXiv.txt": {
        "abstract": "Even though the existence of intermediate-mass black holes (IMBHs, black holes with masses  ranging between $10^{2-4}\\,M_{\\odot}$) has not yet been corroborated observationally, these objects are of high interest for astrophysics. Our understanding of the formation and evolution of supermassive black holes (SMBHs), as well as galaxy evolution modeling and cosmography would dramatically change if an IMBH were to be observed. From a point of view of traditional photon-based astronomy, {which relies on the monitoring of innermost stellar kinematics}, the {\\em direct} detection of an IMBH seems to be rather far in the future. However, the prospect of the detection and characterization of an IMBH has good chances in lower-frequency gravitational-wave (GW) astrophysics using ground-based detectors such as LIGO, Virgo and the future Einstein Telescope (ET). We present an analysis of the signal of a system of a binary of IMBHs (BBH from now onwards) based on a waveform model obtained with numerical relativity simulations coupled with post-Newtonian calculations at the highest available order. IMBH binaries with total masses between $200-20000\\,M_\\odot$ would produce significant signal-to-noise ratios (SNRs) in advanced LIGO and Virgo and the ET. We have computed the expected event rate of IMBH binary coalescences for different configurations of the binary, finding interesting values that depend on the spin of the IMBHs. The prospects for IMBH detection and characterization with ground-based GW observatories would not only provide us with a robust test of general relativity, but would also corroborate the existence of these systems. Such detections should allow astrophysicists to probe the stellar environments of IMBHs and their formation processes. ",
        "introduction": "\\label{sec.motivation}  By following the stellar dynamics at the center of our Galaxy, we have now the most well-established evidence for the existence of a SMBH.  The close examination of the Keplerian orbits of the so-called S-stars (also called S0-stars, where the letter ``S'' stands simply for source) has revealed the nature of the central dark object located at the Galactic Center. By following S2 (S02), the mass of SgrA$^*$ was estimated to be about $3.7\\times 10^6\\,M_{\\odot}$ within a volume with radius no larger than 6.25 light-hours \\citep{SchoedelEtAl03,GhezEtAl03b}. More recent data based on 16 years of observations set the mass of the central SMBH to $\\sim 4 \\times 10^{6} \\, M_{\\odot}$ \\citep{EisenhauerEtAl05,GhezEtAl05,GhezEtAl08,GillessenEtAl09}.  Massive black holes in a lower range of masses may exist in smaller stellar systems such as globular clusters. These are called intermediate-mass black holes (IMBHs) because their masses range between $M\\sim 10^{2}$ and $M\\sim 10^{4}\\,M_\\odot$, assuming that they follow the observed correlations between SMBHs and their host stellar environments. Nevertheless, the existence of IMBHs has never been confirmed, though we have some evidence that could favor their existence \\citep[see][ and references therein]{MillerColbert04,Miller09}.  If we wanted to apply the same technique to detect IMBHs in globular clusters as we use for SMBHs in galactic centers, we would need ultra-precise astronomy, since the sphere of influence of an IMBH is {$\\sim$ few arc seconds. For instance, for a $10^4\\,M_{\\odot}$ IMBH, the influence radius is of $\\sim 5''$ assuming a central velocity dispersion of $\\sigma = 20\\,{\\rm km\\,s}^{-1}$ and a distance of $\\sim 5$ kpc.} The number of stars enclosed in that volume is only a few. Currently, with adaptive optics, one can aspire --optimistically-- to have a couple of measurements of velocities if the target is about $\\sim 5$ kpc away on a time scale of about 10 yrs. The measures depend on a number of factors, such as the required availability of a bright reference star in order to have a good astrometric reference system.  Also, the sensitivity limits correspond to a K-band magnitude of $\\sim$ 15, (B- MS stars at 8 kpc, like e.g. S2 in our Galactic Center).  This means that, in order to detect an IMBH or, at least, a massive dark object in a globular cluster center by following the stellar dynamics around it, one has to have recourse to the Very Large Telescope interferometer and to one of the next-generation instruments, the VSI or GRAVITY \\citep{GillessenEtAl06,EisenhauerEtAl08}. In this case we can hope to improve the astrometric accuracy by a factor of $\\sim 10$. Only in that scenario would we be in the position of following closely the kinematics around a potential IMBH so as to determine its mass.  GW astronomy could contribute to IMBH detection. In the past years, the field has reached a milestone with the construction of an international network of GW interferometers that have achieved or are close to achieving their design sensitivity. Moreover, the first-generation ground-based detectors LIGO and Virgo will undergo major technical upgrades in the next five years that will increase the volume of the observable universe by a factor of a thousand \\footnote{\\scriptsize{http://www.ligo.caltech.edu/advLIGO/, http://wwwcascina.virgo.infn.it/advirgo/}}.  The data that will be taken by the advanced interferometers are expected to transform the field from GW detection to GW astrophysics. The availability of accurate waveform models for the full BBH coalescence in order to construct templates for match-filtering is crucial in the GW searches for compact binaries.  The construction of this kind of templates has recently been made possible thanks to the combination of post-Newtonian calculations of the BBH inspiral and numerical relativity simulations of the merger and ringdown. Two approaches to this problem are the effective-one-body techniques~\\citep{Buonanno:1998gg,Buonanno:2009qa} and the phenomenological matching of PN and NR results~\\citep{Ajith:2007kx,Ajith:2009bn,Santamaria:SpinTemp}. In this paper we use the results of the latter.  The structure of this paper is as follows: In section~\\ref{sec:astro} we expand the astrophysical context to this problem and give a description of the different efforts made to address the evolution of a BBH in a stellar cluster from its birth to its final coalescence.  In section~\\ref{sec:wfmodel} we introduce the techniques used in the data analysis of the waveform modeling of BBH coalescences, present our hybrid waveform model, and use the model to compute and discuss expected signal-to-noise ratios in present and future GW detectors. The use of the new waveform model allows us to give an improved estimate for the  number of events one can expect for several physical configurations in Advanced LIGO and the Einstein Telescope, which we present in section~\\ref{sec:events}. We conclude with a summary of our results and future prospects of our work in section~\\ref{sec:concl}. ",
        "conclusions": "\\label{sec:concl}  Even though we do not have any evidence of IMBHs so far, a number of theoretical works have addressed their formation in dense stellar clusters. If we were to follow the same techniques that have led us to discover the SMBH in our own Galaxy, we would need the Very Large Telescope interferometer and next-generation instruments, such as the VSI or GRAVITY, which should be operative in the next $\\sim 10$ yrs. An alternative or even complementary way of discovering IMBHs is via their emission of GWs when they are in a BBH system.  The identification and characterization of these systems rely on accurate waveform modeling of their GW emission, which has been made possible by the success of numerical relativity in simulating the last orbits of the BBH coalescence and the coupling of these results to analytical post-Newtonian calculations of the inspiral phase. We use a PN-NR hybrid waveform model of the BBH coalescence based on a construction procedure in the frequency domain \\citep[see][for details]{Santamaria:SpinTemp}.  Using this hybrid waveform, we have estimated the SNR corresponding to the current and Advanced LIGO and Virgo detectors, the proposed ET and the space-based LISA at a distance of $z=1$, i.e. 6.68 Gpc. The results indicate that IMBH binaries will produce SNRs sufficient for detection in advanced LIGO and Virgo and notably larger SNRs in the ET, thus making them interesting sources to follow up on. Eventual observations of IMBH binaries with future ground-based detectors could be complementary to those of LISA, which is expected to detect these systems with moderate SNRs and to be more sensitive to SMBH binaries instead. More remarkably, in principle, if LISA and the ET are operative at the same time, they could complement each other and be used to track a particular event.  Furthermore, we have revisited the event rate of BBHs for various detectors and find encouraging results, within the inherent uncertainties of the approach. Our estimations are consistent with previous works, and additionally we have quantified the effect of the total spin of the binary in the expected event rates. We have estimated the distance to which Advanced LIGO and the ET will be able to see binaries of IMBHs. This quantity depends strongly on the mass ratio and spins of the binary. For Advanced LIGO, equal-mass, non-spinning configurations of observed total mass $\\sim 200$--$700\\,M_\\odot$ can be seen up to $z \\sim 0.8$. If the spins are aligned with the total angular momentum and significant ($\\chi_{1,2} \\sim 0.75$), the reach increases to $z \\sim 1.5$ for observed total masses of $\\sim 400\\,M_\\odot$. The ET will be able to explore even more remote distances, reaching to $z\\sim 5$ and further. Our present knowledge of stellar formation at such large redshifts is incomplete, therefore we have computed the event rates for the ET integrating only until $z=5$. We have compared three star formation models and three different configurations of the binary. The effect of the particular formation model is not very significant; remarkably, the physical configuration does however influence the final rates. We provide the results for three physical configurations, taking into account both single- and double-cluster channels in the binary formation. For a fully correct calculation of the event rate integrated over all possible configurations, more detailed knowledge of the distribution of spins and mass ratios of IMBH binaries formed in globular cluster is required.  Advanced ground-based detectors are designed to be able to operate in different modes so that their sensitivity can be tuned to various kinds of astrophysical objects. Considering the importance of an eventual detection of a BBH, the design of an optimized Advanced LIGO configuration for systems with $M \\sim 10^{2-4}\\,M_{\\odot}$ would be desirable to increase the possibility of observing such a system.  In case an IMBH binary coalescence were detected, the recovery and study of the physical parameters of the system could serve to test general relativity and prove or reject alternative theories, such as scalar-tensor type or massive graviton theories. The {\\em direct} identification of an IMBH with GWs will be a revolutionary event not only due to the uncertainty that surrounds their existence and their potential role in testing general relativity. The information encoded in the detection will provide us with a detailed description of the environment of the BBH/IMBH itself."
    },
    "0910/0910.5186_arXiv.txt": {
        "abstract": "%  We report a chromospheric jet lasting for more than 1~hr observed by {\\it Hinode} Solar Optical Telescope in unprecedented detail. The ejection occurred in three episodes separated by 12--14~min, with the amount and velocity of material decreasing with time. The upward velocities range from 438 to $33 \\kmps$, while the downward velocities of the material falling back have smaller values (mean: $-56 \\kmps$) and a narrower distribution (standard deviation: $14 \\kmps$). The average acceleration inferred from parabolic space-time tracks is $141 \\m \\, {\\rm s^{-2}}$, a fraction of the solar gravitational acceleration. The jet consists of fine threads ($0\\farcs5$--$2\\arcsec$ wide), which exhibit coherent, oscillatory transverse motions perpendicular to the jet axis and about a common equilibrium position. These motions propagate upward along the jet, with the maximum phase speed of $744 \\pm 11 \\kmps$ at the leading front of the jet. The transverse oscillation velocities range from $151$ to $26 \\kmps$, amplitudes from $6.0$ to $1.9 \\Mm$, and periods from $250$ to $536 \\s$. The oscillations slow down with time and cease when the material starts to fall back. The falling material travels along almost straight lines in the original direction of ascent, showing no transverse motions. These observations are consistent with the scenario that the jet involves untwisting helical threads, which rotate about the axis of a single large cylinder and shed magnetic helicity into the upper atmosphere. ",
        "introduction": "\\label{sect_intro}  Transient, small-scale ejections of plasma from the lower atmosphere are common manifestations of solar activity. Cool plasma ejections were observed as emission in \\Ha or absorption at other wavelengths and historically called {\\it surges} \\citep{NewtonH.surge-discovery.1934MNRAS..94..472N}. Hot ejections were usually called {\\it jets} and observed as emission in ultraviolet \\citep{Brueckner.Bartoe.UVjet1983ApJ...272..329B}, extreme ultraviolet \\citep[EUV;][]{Alexander.Fletcher.jet.1999SoPh..190..167A}, and soft X-rays \\citep{ShibataK.1st-SXT-jet.1992PASJ...44L.173S, StrongK.SXT-jet1992PASJ...44L.161S}. Some observations \\citep{SchmiederB.cool-surge.hot-jet1988A&A...201..327S, ChaeQiu.jet.surge1999ApJ...513L..75C, JiangYC.surge.jet.twist2007A&A...469..331J} have indicated that cool surges and hot jets are closely related in space and time, but their physical relationship has not been established.\t% Torsional motions or helical features have long been seen in surges or jets \\citep[e.g.,][]{XuAA.surge-rotate.1984AcASn..25..119X, Kurokawa.untwist-filamt1987SoPh..108..251K, ShimojoM.jet-stat.1996PASJ...48..123S, Patsourakos.EUVI-jet2008ApJ...680L..73P}, while their exact causes remain unclear. From now on, we refer to both surges and jets with a general term {\\it jets}, unless otherwise noted.  Solar jets are commonly associated with\t\t% flux emergence \\citep{Roy.surge.magnetic.1973SoPh...28...95R}\t% or moving magnetic features \\citep{Gaizauskas.movingBfield.surge1982AdSpR...2...11G, BrooksDH.surge2007ApJ...656.1197B}. Models involving magnetic reconnection have thus been proposed \t% \\citep[e.g.,][]{RustD.surge-reconn.1968IAUS...35...77R, YokoyamShibata.jetModel1995Natur.375...42Y}. Mechanisms \\citep[see review in][]{Canfield.surge-jet1996ApJ...464.1016C}\t\t% suggested for accelerating jet material to 10--1000~$\\kmps$ include reconnection outflows driven by the slingshot effect of magnetic tension \\citep{YokoyamShibata.jetModel1995Natur.375...42Y}, the pressure gradient behind the shock formed by reconnection outflows \\citep{YokoyamShibata.jetModel1996PASJ...48..353Y, Tarbell.Ryutova.jet.1999ApJ...514L..47T}, chromospheric evaporation caused by heating from the associated flare \\citep{ShimojoM.jet-phys.2000ApJ...542.1100S}, and relaxation of magnetic twists \\citep{Shibata.Uchida.helic-jetMHD.1985PASJ...37...31S}.   \\hinode \\citep{Kosugi.Hinode2007SoPh..243....3K}, with its superior resolutions, has provided unprecedented details and spurred renewed interest in solar jets. They were found to be ubiquitous on various spatial and temporal scales in the \\ion{Ca}{2} H line \\citep{Shibata.CaHjet.2007Sci...318.1591S, Nishizuka.giantCaHjet.2008ApJ...683L..83N}, EUV \\citep{Culhane.EIS-jet2007PASJ...59S.751C, Moreno-Insertis.EISjet2008ApJ...673L.211M}, and soft X-rays \\citep{Savcheva.XRT.jet.stat2007PASJ...59S.771S, NittaN.jetHe3.2008ApJ...675L.125N}. In this Letter, we report an intriguing jet observed by \\hinode in great detail, and attempt to obtain new clues to some unanswered questions mentioned above. \\begin{figure*}[thb]      % \\epsscale{1.2}\t% \\plotone{f1.eps} \\caption[]{ Negative \\hinode \\ion{Ca}{2} H images.\t\t% The diagonal dotted line \t% marks the jet axis. The dashed box indicates a $10\\arcsec$ wide cut along the axis for the time-distance diagram shown in Figure~\\ref{para.eps}(a),\t\t% while the numbered dotted lines represent $1\\arcsec$ narrow cuts perpendicular to the axis (see online movie~1). } \\label{maps.eps} \\end{figure*} ",
        "conclusions": "% \\label{sect_conclude}  We have presented kinematic measurements of both axial and transverse motions of a chromospheric jet observed by \\hinode SOT at high spatial and temporal resolutions. This study complements previous low resolution counterparts of \\Ha surges \t% and EUV or X-ray jets, and offers new insights to this type of phenomena. Our major results and interpretations are as follows.  \\begin{enumerate}  \\item\t% The ejection occurs in three episodes separated by 12--14~min, rather than continuously. The amount and velocity of ejected material decrease with time.\t% The ejecting velocities have a wide range from 438 to $33 \\kmps$, while the velocities of material falling back \t\t% have a narrow range (mean: $-56 \\kmps$ and standard deviation: $14 \\kmps$). The acceleration inferred from parabolic tracks in the time-distance diagram has a mean of $141 \\m \\, {\\rm s^{-2}}$, a fraction of the solar gravitational acceleration.  \\item\t% The jet consists of $0\\farcs5$--$2 \\arcsec$ thick fine threads, which exhibit oscillatory transverse motions across the jet about a common equilibrium position. These oscillations have velocities ranging from $151 \\pm 6$ to $26 \\pm 3 \\kmps$, amplitudes from $6.0 \\pm 0.2$  to $1.9 \\pm 0.2 \\Mm$, and periods from $250 \\pm 6$ to $536 \\pm 19 \\s$. The upward propagation of the oscillations has a maximum phase speed of $744 \\pm 11 \\kmps$ (comparable to the coronal Alfv\\'{e}n speed) associated with the leading front of the jet.  \\item\t% The transverse motions slow down with time and cease near the time when the material reaches its maximum height and starts to fall back. The falling material travels along almost straight lines in the original direction of ascent, showing no more signatures of transverse motions.  \\item\t% These observations are consistent with the scenario that the jet involves unwinding of left-handed helical threads that rotate counter-clockwise about a common axis. The untwisting wave front propagates upward at the Alfv\\'{e}n speed. The pitch spacing between adjacent helical threads increases with height, consistent with the expected axial expansion of\t% unwinding screws. The jet results in\t% magnetic helicity being shed into the upper atmosphere.   \\end{enumerate}  A more in-depth multiwavelength study of this event and comparison with theoretical models are underway and will be published in the future."
    },
    "0910/0910.5715_arXiv.txt": {
        "abstract": "{Clusters of galaxies are believed to be capable to accelerate  protons  at accretion shocks to  energies exceeding   $10^{18}$ eV. At these energies, the losses caused by  interactions of cosmic rays  with  photons of the  Cosmic Microwave Background Radiation (CMBR)  become effective and determine the maximum energy of protons and the shape of the energy spectrum  in the cutoff region.} {The aim of this work is the  study of the formation of the energy spectrum of accelerated protons at  accretion shocks of galaxy clusters and of the characteristics of  their broad band emission.} {The proton  energy distribution is calculated  self-consistently via a time-dependent numerical treatment of the shock acceleration process which takes into account the proton energy losses due to interactions with the CMBR.} {We calculate the energy distribution of accelerated protons, as well as the flux  of broad-band emission produced  by secondary electrons and positrons via synchrotron and inverse Compton scattering processes. We find that the downstream and upstream regions contribute almost at the same level to the emission. For the typical parameters characterising galaxy clusters, the synchrotron and IC peaks in the  spectral energy distributions appear  at comparable flux levels.} {For an efficient acceleration, the expected emission components in  the X-ray and gamma-ray band  are close to the detection threshold of  current  generation instruments, and will be possibly detected with the future generation of detectors.} ",
        "introduction": "\\label{Introduction}  Rich clusters of galaxies are the largest virialised structures in the Universe, with typical sizes of a few Mpc and masses up to $10^{15} M_{\\odot}$ or more (see \\citealp{sarazinbook} for a review). In the standard picture of cosmic structure formation, the structure's growth is driven by gravitational instability. This process is hierarchical, with larger systems forming later via the assembly of pre-existing smaller structures. Within this scenario, galaxy clusters form via mergers, and their age can be estimated to be of the order of 10 Gyr (e.g. \\citealp{agecl}). In addition, cold material from the surrounding environment is continuously infalling, due to gravitational attraction, and an expanding shock wave, called the accretion shock, is expected to form at the cluster boundary and to carry outward the information of virialisation \\citep{bertschinger}. Numerical simulations have confirmed the appearance of the so--called accretion shocks during structure formation (e.g. \\citealp{kangsim}).  The detection of a tenuous and diffuse synchrotron radio emission from about one third of rich clusters of galaxies (e.g. \\citealp{govoni}) reveals the presence of a diffuse magnetic field and a population of high energy electrons. Using exclusively synchrotron data, though, it is possible to get information only on the product of particles density and magnetic field energy density (unless adopting assumptions like equipartition). Detection of non-thermal X-rays have also been claimed from a few of such sources and generally interpreted as inverse Compton emission from the same population of electrons (e.g. \\citealp{ff, eckert}). Since IC depends on the electron density but not on the magnetic field, combining the data allows us to break the degeneracy and thus determine values of B of the order of $0.1 \\mu$G. Significantly higher values, of the order of a few microGauss, are obtained from Faraday rotation measurements \\citep{Clarke,carilli}. Thus the uncertainty in the determination of the magnetic field is quite large, allowing values to vary of about one order of magnitude \\citep{nnrephaeli}.  According to theoretical models, the electrons responsible for the radio and non-thermal X-ray emission, are produced through different acceleration mechanisms (e.g. \\citealp{schlickeiser,tribble,sarazin,brunetti2001,petrosian2001}). Alternatively, the radio synchrotron radiation  can be emitted by secondary electrons produced at interactions of accelerated protons with the intracluster gas \\citep{dennison,blasicol,atoyan,dolag,kushnir}. In this scenario, gamma ray emission is expected due to the decay of neutral pions produced in \\textit{p-p} interactions \\citep{volk_cluster,ber_cluster}. The non-thermal X-rays emission also can be related to synchrotron radiation of secondary electrons, but of much higher energies, produced in photon-photon  \\citep{timokhin} and proton-photon \\citep{FA02,susumu} interactions.  Several particle acceleration mechanisms have been proposed to operate in clusters of galaxies (see \\citealp{revste} for a review). In particular, it has been argued that large scale shocks can effectively accelerate electrons and protons up to ultrarelativistic energies \\citep{norman,ber_cluster,loeb,miniati,blasimerger,steshocks,ryu,berrington,pfrommer}. In a recent paper we performed  detailed calculations  of diffusive shock acceleration of electrons in galaxy clusters \\citep{mio}. Electrons can be accelerated up to 100 TeV at cluster accretion shocks,  with synchrotron X-ray and  IC gamma-ray radiation components produced mainly in the downstream region. While the maximum energy of electrons is limited by synchrotron and IC energy losses, the protons can be accelerated to much higher energies.  Indeed, according to the Hillas criterion \\citep{hillas},  galaxy clusters are  amongst the few source populations capable, as long as this concerns the dimensions of the structure and the value of the magnetic field, to accelerate protons up to $10^{20}$ eV.  Moreover, clusters are cosmological structures and their lifetime is comparable to the age of the Universe, therefore, if acceleration takes place in such objects, it can continue up to $\\sim 10^{10}$ yr. On the other hand, high energy protons up to energies of the order of $10^{15}$ eV and possibly higher are well confined in the volume of the cluster over this time scale \\citep{volk_cluster,ber_cluster}. This results in an effective accumulation of high energy particles in the cluster. The thermal energy budget of rich clusters, estimated from measurements of their thermal X-ray emission, is of the order of $10^{63}-10^{64}$ erg. The non-thermal component seems to be constrained to a few percents of the thermal energy $\\sim 10^{62}$ erg \\citep{veritas,coma,magic}. This value is also comparable with the magnetic field energy, assuming 1 $\\mu$G and a spherical cluster of 3 Mpc of radius (e.g. the size of the Coma cluster).  Although the dimensions of the system, the strength of the magnetic field,  and the age of the accelerator formally allow protons to be accelerated up to $10^{20}$ eV, the particles lose, in fact,  their energy via pair and pion production in the interactions with the photons of the Cosmic Microwave Background (CMBR) radiation field. As discussed in \\citet{norman} and \\citet{kang_rach_bier}, for a shock velocity of a few thousands km/s and a magnetic field of the order of  $1 \\mu \\rm G$, the shock acceleration rate is compensated by the energy loss rate at energies around $10^{19}$ eV (the exact value depends on the assumed diffusion coefficient and also on the shock geometry as shown by \\citealp{ostrowski}).  The interactions of ultra-high energy protons with the CMBR lead to production of electrons in the energy domain ($\\geq 10^{15} \\ \\rm eV$) which is not  accessible through any direct acceleration mechanism. These electrons   cool  via synchrotron radiation and IC scattering   on very short  timescales (compared to  both the age of the source and  interaction timescales of protons).  For the same reason they are rather  localised in space, i.e. are ``burned\" not far from the sites of their production. Therefore the corresponding radiation components in the X-ray and gamma-ray energy bands are  precise tracers of primary protons, containing information about the acceleration and propagation of their ``grandparents\". This interesting  feature has been indicated in  \\citet{FA02} in the general context of  acceleration and propagation of ultrahigh energy protons in large scale extragalactic structures. More specifically, this issue was discussed by \\citet{susumu} for objects like the Coma galaxy cluster. However in these  studies  an {\\it a priori} spectrum of protons has been assumed in the ``standard form\"  $E^{-2} \\exp(-E/E^*)$. In fact,  the  energy distribution of protons accelerated in galaxy clusters via Diffusive Shock Acceleration (DSA) mechanism  can be quite different from this simple form, as it is demonstrated below.  In this paper we  study the process of proton acceleration by accretion shocks in galaxy clusters taking into account self-consistently the  energy loss channels of protons related to their interactions with  the CMBR photons. We make use of the numerical approach presented in \\citet{mio}. In Section \\ref{Proton Acceleration and Energy losses} the calculation is introduced and the accelerated proton spectra are derived. In particular, we show that electron/positron pair production is the dominant energy loss channel for protons. In Section \\ref{Spectra of Secondary Electrons} we calculate the spectra of  secondary pairs. The broad-band emission produced by the secondary electrons during their interactions with the background magnetic and radiation fields is presented in Section \\ref{Radiation Spectra}. In Section \\ref{Shock Modification} we study, in a simplified fashion, the impact of shock modification by efficiently accelerated protons on the acceleration and emission features.  We  briefly discuss and summarise the main results  in Section \\ref{Conclusions}. ",
        "conclusions": "\\label{Conclusions}  Proton acceleration in galaxy clusters was studied in the framework of DSA via a detailed time-dependent numerical calculation that includes energy losses due to interactions of protons with photons of the CMBR. For realistic  shock speeds of a few thousand km/s and a background magnetic field close to  $1 \\mu  \\rm G$, the maximum energy achievable by protons is determined by the energy losses  due to  pair production and ranges from a few times $10^{18}$ eV to a few times $10^{19}$ eV.  We performed the calculations assuming that acceleration takes place on time scales comparable to the age of the cluster. Since steady state is never achieved in this scenario, a time-dependent treatment is required. Particle spectra, when calculated including the effect of energy losses, exhibit interesting features. The decay of the spectrum above the cut-off energy is not exponential. Its dependence on energy is shallower due to the flat profile of pair production timescales in the cut-off energy range. The time-dependent distributions of protons are used to calculate accurately the  production rates of secondary electron-positron pairs. These electrons cool rapidly via synchrotron radiation and IC scattering which proceeds in the Klein-Nishina  regime. The effect of the hardening induced by the KN cross-section is visible in the IC radiation spectra both in the upstream and downstream regions of the shock. For the fiducial Coma-like cluster used in this work, the synchrotron and IC peaks of the electron broadband SED  are at comparable levels and the associated flux from a source at the distance of a 100 Mpc is expected at the level of $10^{-12}$ erg cm$^{-2}$ s$^{-1}$ in the X-rays and an order of magnitude lower for TeV gamma-rays. Note that  although the maximum of the gamma-ray emission is located above 100 TeV, it unfortunately cannot be observed due to severe  intergalactic absorption. The expected  gamma-ray flux  from clusters of galaxies is at the limit of the sensitivity of present generation instruments, however it may be detectable with the future generation of detectors. The optimum energy interval for  gamma-ray detection  is between 1 and 10 TeV.  The  detectability of clusters in hard X-rays and  gamma-rays  associated with  interactions of ultrahigh energy protons with the CMBR, depends on the value of the parameter $A=W_{62}/d_{100}^2$, where $W_{62}=W/10^{62} \\ \\rm erg$ is the total energy released in cosmic rays normalised to $10^{62}$ erg, and $d_{100}=d/100$ Mpc is the distance normalised to 100 Mpc. Obviously, the best candidates  for detection are nearby  rich galaxy clusters like Coma and Perseus located at distances $d \\sim 100$ Mpc. On the other hand, due to the large extension of the  non-thermal  emission  (as large as several Mpc), and given that the angular size of the source $\\theta \\propto 1/d$,  the  probability of detection reduces  with the distance slower than  $1/d^2$. Nevertheless, as long as the total energy in accelerated protons  does not significantly exceed $10^{62}$ erg, the visibility of  clusters of galaxies in X-rays and gamma-rays is limited by objects located within a few 100 Mpc.  The chances of detection of non-thermal emission of clusters related to ultrahigh energy protons, especially in the hard X-ray band, can be significantly higher if protons are accelerated by  non-linear shocks modified by the pressure of relativistic particles. In this scenario, a large fraction of the energy of the shock is transferred to relativistic protons. Moreover, in this case a very hard spectrum of protons is formed, thus the main fraction of non-thermal energy is carried by the highest energy particles.  These two factors can enhance the luminosity of X- and gamma-ray emission of secondary electrons by more than an order of magnitude, and thus  increase the probability of detection of clusters located beyond 100 Mpc.  In this paper we do not discuss the gamma-ray production related to interactions of accelerated protons with the ambient gas which can compete with the inverse Compton radiation of pair produced electrons. The relative contributions of these two channels depends on the density of the ambient gas and the spectral shape of accelerated protons. The flux of gamma-rays from {\\it pp} interactions can be easily  estimated based on the cooling time of protons, $t_{\\rm pp} \\approx 1.5 \\times 10^{19}/n_{-4}$ s, where $n_{-4}=n/10^{-4}$ cm$^{-3}$ is the density of the ambient hydrogen gas, normalised to $10^{-4}$ cm$^{-3}$.  Then, the energy flux of gamma-rays at 1 TeV is estimated as $F_\\gamma (\\sim 1 \\rm TeV) \\approx  6 \\times 10^{-12} \\kappa W_{62} n_{-4}/d_{100}$ erg cm$^{-2}$ s$^{-1}$, where $\\kappa$ is the fraction of the total energy of accelerated protons in the energy interval between 10 to 100 TeV  (these protons are primarily responsible for production of gamma-rays of energy $\\sim 1$ TeV). For a proton energy spectrum extending to $10^{18}$ eV, this fraction is of order of $\\kappa \\sim 0.1$. Thus for an average gas density in a cluster like Coma, $n \\sim 3 \\times 10^{-4}$ cm$^{-3}$, the gamma-ray flux at 1 TeV is expected at the level of $10^{-12}$ erg cm$^{-2}$ s$^{-1}$ which is comparable to the contribution of IC radiation of secondary electrons. In the case of harder spectra of protons accelerated by non-linear shocks, the contribution of gamma-rays from {\\it pp} interactions is dramatically reduced and the contribution of  secondary pairs to gamma-rays via IC scattering  strongly dominates over gamma-rays from {\\it pp} interactions."
    },
    "0910/0910.1136_arXiv.txt": {
        "abstract": "{OH maser emission at 1.67 GHz is known to be associated with regions of intense star formation within ULIRGs. As these galaxies are formed through violent mergers, studying the co-moving density of OH maser galaxies across cosmic time will allow the merger rate of the Universe to be determined in an independent way. This merger rate is an important parameter in galaxy formation and evolution scenarios. The sensitivity, wide field of view and spectral coverage of both APERTIF on the WSRT and ASKAP will allow for the first time all-sky blind surveys for OH maser galaxies to be carried out to redshift 1.4. We describe the prospects for such surveys, including the expected number of OH maser galaxies that will be discovered, and what limits can be placed on the OH maser luminosity function, and hence the merger rate out to redshift 1.4 with various survey strategies.}  \\FullConference{Panoramic Radio Astronomy: Wide-field 1-2 GHz research on galaxy evolution - PRA2009\\\\ June 02 - 05 2009\\\\ Groningen, the Netherlands}  \\begin{document} ",
        "introduction": "The 1.665 GHz and 1.667 GHz doublet transitions of hydroxyl (OH) show spectacularly luminous maser emission in regions of intense circumnuclear star-formation, with reported line luminosities of up to $10^4~L_{\\odot}$ (for a review see \\cite{lo05}; for an example of the OH doublet maser lines see Figure \\ref{maser-red}). The masing action results from a strong far-infrared radiation field, generated by the heated dust associated with star-formation, that pumps the OH molecules into an excited state \\cite{baan85,baan89}. The most powerful OH masers are found in the population of ultra-luminous infrared galaxies (ULIRGs; $L_{FIR} \\geq 10^{12}~L_{\\odot}$) that are believed to form through galaxy mergers. Indeed, surveys of ULIRGs at low redshifts have found up to one third to show OH maser emission \\cite{darling02a}, with a correlation found between the far-infrared and OH maser luminosities ($L_{OH} \\propto L_{FIR}$$^{2.3}$; \\cite{kloeckner04}). Thus, one of the most powerful applications of OH masers is their unbiased ability to trace the number density of ULIRGs and hence the galaxy merger history of the Universe, which forms a fundamental part of galaxy formation and evolution scenarios (e.g. \\cite{croton06}). There are to date $\\sim$100 OH maser galaxies known and all are found in the local Universe ($z \\leq 0.3$; see Figure \\ref{maser-red} for the OH maser redshift distribution).  By detecting OH masers at higher redshifts, it will be possible to determine the density evolution of the OH maser luminosity function, and hence make an independent measurement of the merger rate of the Universe as a function of redshift.  In the future, wide-field surveys with the next generation of radio telescopes and arrays can be carried out to determine the evolution in the OH maser luminosity function. The potential for using OH maser galaxies in this way was pointed out by Briggs \\cite{briggs98}.  In this paper, we discuss the prospects for such surveys with two of these new facilities; the first is APERTIF (APERture Tile In Focus) on the Westerbork Synthesis Radio Telescope and the second is ASKAP (the Australian Square Kilometre Array Pathfinder).  \\begin{figure}[tbh] \\begin{center} \\includegraphics[width=\\textwidth]{maser-red.eps} \\caption{Left: an example of the 1.665 GHz and 1.667 GHz doublet transitions of hydroxyl from the galaxy 10035+2740 at redshift 0.1662, taken from \\cite{darling02a}. Right: the redshift distribution of extragalactic OH maser galaxies. Note that this redshift distribution is not complete, but is based on targeted observations of candidate OH maser galaxies, for example, ULIRGs.} \\label{maser-red} \\end{center} \\end{figure} ",
        "conclusions": "\\label{conc}  We have presented new simulations of the potential numbers of OH maser galaxies that can be detected with APERTIF and ASKAP out to redshift $\\sim$1.4. Our results suggest that with 3 months of observing time, it will be possible to cleanly differentiate between evolutionary models for the OH maser population. Given that these new facilities will be carrying out large all-sky blind surveys for H\\,{\\sc i} in the same frequency range that we would expect to find high redshift OH, it would require only a small additional amount of effort to expand the search to look for OH galaxies.  The largest uncertainty in our analysis comes from the poorly constrained local OH maser luminosity function at high luminosities ($\\geq 10^{3.8}~L_{\\odot}$). There are a number of avenues available to try and constrain the high-luminosity end of the luminosity function in the short term, which would further constrain the predictions made here. The first is to carry out blind searches of small fields with current instruments, for example, with the Giant Metrewave Radio Telescope, as suggested by Darling \\& Giovanelli \\cite{darling02b}. Another option is to observe gravitationally lensed ULIRGs to detect OH masers at even higher redshifts ($z>2.5$). Here, the lensing magnification can be used to detect faint lines which would otherwise not be observable without the magnification of the lens. This technique was recently used to detect the most distant water maser from a high redshift galaxy \\cite{impellizzeri08}."
    },
    "0910/0910.4978_arXiv.txt": {
        "abstract": "{} {It is widely recognized that processes taking place inside group environment could be among the main drivers of galaxy evolution. Compact groups of galaxies are in particular good laboratories for studying galaxy interactions and their effects on the evolution of galaxies due to their high density and low velocity dispersion. SCG0018-4854 is a remarkably high galaxy density and low velocity dispersion group with evidence of a recent interaction. It has been detected and analyzed at different wavelengths, but its kinematics has not yet been studied in detail.} {We obtained VLT FORS2 optical observations and we present spectroscopic and photometric evidence of how dramatically galaxy interactions have affected each of the four member galaxies.} {We found peculiar kinematics for each galaxy and evidence of recent star formation. In particular, the gas and stellar radial velocity curves of two galaxies are irregular with a level of asymmetry similar to that of other interacting galaxies. We discovered the presence of a bar for NGC 92 therefore revising a previous morphological classification and we obtained spectroscopic confirmation of a galactic-scale outflow of NGC 89.} {Peculiar kinematics and dynamic consideration lead to a rough estimate of the age of the latest interaction: $\\tau \\approx 0.2-0.7 \\, \\mathrm{Gyr}$, suggesting that SCG0018-4854 is a young and dynamical group.} ",
        "introduction": "Compact groups (CGs) are among of the smallest and densest systems of galaxies of the universe, furthermore they are characterized by low velocity dispersions of the order of $\\approx 200\\, \\mathrm{km \\, s^{-1}}$ \\citep{1992ApJ...399..353H}. This is why they are excellent laboratories for studying galaxy interactions and their effects on the evolution of galaxies. \\citet{1991ApJS...76..153R} showed that 30\\% of spiral galaxies in CGs has distorted rotation curves and another 30\\% has irregular velocity patterns indicating that tidal interactions are frequent and persisting in CGs. Moreover, \\citet{1994ApJ...427..684M} demonstrated that 43\\% of the galaxies in their compact group sample shows morphological and/or kinematical distortions indicative of interactions and/or mergers. Since these pioneering studies it has been clear that interactions play an important role in the evolution of groups. According to current dynamical models, galaxies belonging to groups should interact violently and merge on a relatively short time scale \\citep{1987ApJ...322..605S,1985MNRAS.215..517B,1996ApJ...458...18G}. Yet, strong mergers in CGs are rare \\citep{1994ApJ...427..684M} and a lot of effort has been made to establish the origin and fate of such systems. Recent studies \\citep{2007AJ....133.2630C} have shown that galaxies in compact groups are more likely to merge under dry conditions, after they have lost most of their gas in interactions among galaxies. These hypothesis is corroborated by HI studies; \\citet{2001A&A...377..812V} have found out perturbed distributions of neutral hydrogen in CGs. This results indicate that the evolution of CGs is mainly dominated by galaxy interactions, continuous tidal stripping and/or gas heating.  It is therefore interesting to reconstruct the evolution and the interaction history of a CG throughout the properties of its galaxy members and compare them with group properties as a whole. Kinematical and dynamical information are fundamental to separate between different interaction histories. \\citet{1998ApJ...507..691M} defined some kinematic indicators to distinguish between merging and interaction, such as misalignment between the kinematic and photometric axis of gas and stellar components, double gas components, warping and other peculiarities. However, there are still open questions about how many interactions take place and how strong and frequent they are.  In this work we focus on SCG0018-4854, a spiral-only compact group characterized by a really dense environment in which member galaxies have clearly interacted.  The plan for the article is the following. In Sect.~\\ref{sec:SCG0018-4854}, we present the characteristics of SCG0018-4854. In Sect.~\\ref{sec:obs}, we describe our observations and the reduction process. In Sect.~\\ref{sec:method}, we explain the methods used for our analysis. The results are presented in Sect.~\\ref{sec:results}, followed, in Sect.~\\ref{sec:Discussion}, by a short discussion and our conclusions in Sect.~\\ref{sec:Conclusions}. ",
        "conclusions": "\\label{sec:Conclusions}  We have analyzed high signal-to-noise spectra along both the major and the minor axis of each galaxy of SCG0018-4854. Each galaxy shows disturbed or peculiar kinematics. Kinematic information is important in order to have some insight about the possible interaction history of the galaxies. Different interaction scenarios, depending on the strength and the geometry of the encounter and the morphological types of the interacting systems, will leave different signatures on the galaxy kinematics. Following \\citet{1998ApJ...507..691M} we found out signs of disturbed velocity fields, asymmetric rotation curves, multiple kinematic gas components, tidal tails and nuclear activity. All these characteristics suggest that at least three out of four galaxies of SCG0018-4854 have been involved in a strong and recent interaction. We derived some constraints for the age of the latest close interaction: $\\approx 190 \\, \\mathrm{Myr} < \\mathit{\\tau_{coll}} < 0.7 \\, \\mathrm{Gyr}$. The interaction has been strong enough to perturb the gas kinematics up to a level of 24\\% for the gas of the two main galaxies. These results are in agreement with the recent claim by \\citet{2009arXiv0908.2798T} that this group is in an advanced stage of interaction, based on the presence of young star forming regions and a candidate tidal dwarf galaxy identified in the UV in addition to the observed tidal tails. Finally, the results of our spectroscpic analysis are consistent with the hypothesis of a large-scale outflow in NGC 89.  In contrast to its evolved properties as a group such as the high spatial density, the spatial distribution of HI and the presence of a common hot IGM, SCG0018-4854 is entirely composed of spiral galaxies and has a remarkably high fraction of active galaxies. Moreover, this study reveals kinematical signs of recent interactions among its members making SCG0018-4854 a dynamically young group although its global properties suggest a more advanced stage of evolution.  Given the estimated age of the latest interaction, we could say that we catch these galaxies in the act of interacting. What about the future of this group? Hydrodynamical simulations could help us to understand the fate of SCG 0018-4854, a really isolated interacting compact group."
    },
    "0910/0910.1593_arXiv.txt": {
        "abstract": "Merging compact binaries are the one source of gravitational radiation so far identified. Because short-period systems which will merge in less than a Hubble time have already been observed as binary pulsars, they are important both as gravitational wave sources for observatories such as LIGO but also as progenitors for short gamma-ray bursts (SGRBs). The fact that these systems must have large systemic velocities implies that by the time they merge, they will be far from their formation site. The locations of merging sites depend sensitively on the gravitational potential of the galaxy host, which until now has been assumed to be static. Here we refine such calculations to incorporate the temporal evolution of the host's gravitational potential as well as that of its nearby neighbors using cosmological simulations of structure formation. This results in merger site distributions that are more diffusively distributed with respect to their putative hosts, with locations extending out to distances of a few Mpc for lighter halos. The degree of mixing between neighboring compact binary populations computed in this way is severely enhanced in environments with a high number density of galaxies. We find that SGRB redshift estimates based solely on the nearest galaxy in projection can be very inaccurate, if progenitor systems inhere large systematic kicks at birth. ",
        "introduction": "The association of short gamma-ray bursts with both star-forming galaxies and with ellipticals dominated by old stellar populations \\citep{Berger05,Bloom05,Fox05,Gehrels05,Prochaska05,Berger09} suggested an analogy to type Ia supernovae, as it indicated a class of progenitors with a wide distribution of delay times between formation and explosion. Similarly, just as core-collapse supernovae are discovered almost exclusively in star-forming galaxies, so too are long GRBs \\citep{wbrev}. Indeed, a detailed census of the types of host galaxies, burst locations and redshifts could help decide between the various SGRB progenitor alternatives \\citep[e.g.][]{zr07,guetta05,bp06}: If the progenitor lifetime is long and the systemic kick is small, then the bursts should correspond spatially to the oldest populations in a given host galaxy. For early-type galaxies, the distribution would most likely follow the light of the host. A neutron star (NS) binary could take billions of years to spiral together, and could by then, if given a substantial kick velocity on formation, merge far from its birth site \\citep{fryer99,Bloom99,bbk02,Bloom02}.  The burst offsets would then most likely be larger for smaller mass hosts.  Double neutron star binaries, such as the famous PSR1913+16, will eventually coalesce, when gravitational radiation drives them together \\citep{vk07}. Each supernova is thought to impart a substantial kick to the resulting NS \\citep{hp97}.  For systems that survive both supernovae explosions the center of mass of the remnant binary itself will receive a velocity boost on the order of a few hundred kilometers per second \\citep{bp95,fk97}.  As a result, NS binaries will be ejected from their birth sites.  The exact distribution of merger sites depends sensitively on the gravitational potential of the host, which until now has been assumed to be static. The potential of a realistic host galaxy is, however, not static. In fact, the gravitational potential of the host as well as that of its nearby neighbors is expected to evolve dramatically from compact binary production until coalescence.  In order to incorporate these effects self-consistently, in this {\\it Letter}, we study the orbital evolution of compact binary systems using cosmological simulations of structure formation.  Our results provide new insights into what happens when compact binary stars are ejected from their birth halos as a result of velocity kicks, and what progenitor clues a distant observer might uncover from the distribution of SGRB sites in and around galaxies. ",
        "conclusions": "Before the detailed modeling of light curves was used to constrain the nature of supernovae progenitors, the location of supernovae in and around galaxies provided important clues to the nature of the progenitors.  Similarly, in the absence of supernova-like features \\citep[e.g.][]{hjorth05a}, detailed observations of the astrophysics of individual host galaxies may thus be essential before stringent constraints on the identity of SGRB progenitors can be placed\\footnote{The interaction of burst ejecta with a stellar binary companion \\citep{andrew} or with its emitted radiation \\citep{enrico} could also help shedding light on the identity of the progenitor system.} \\citep[e.g.][]{zr07,Fox05}.  Even with a handful of SGRBs detected to date, it has become apparent that short and long events are not drawn from the same parent stellar population \\citep{nakar}. In contrast to long GRBs, the galaxies associated with SGRBs exhibit a wide range of star-formation rates, morphologies and metallicities \\citep{Berger09}. They are also often found in older and lower-redshift galaxies and, in a few cases, with large ($\\gtrsim$ 10 kpc) projected offsets from the centers of their putative host galaxies \\citep{Bloom05,Prochaska05,Berger05,Gorosabel06,levan,bloom07}.  SGRBs are, however, not universally at large offsets and are not always associated with early-type galaxies \\citep{Fox05,villa,hjorth05b,covino,as06}.  The discovery of early-type galaxy hosts suggest a progenitor lifetime distribution extending well beyond a Gyr. A large progenitor lifetime would help explain the apparent high incidence of galaxy cluster membership \\citep{Bloom05,Ped05,Berger07a}. On the other hand, shorter lifetimes are required to explain the population of SGRB at moderately high redshift \\citep{Berger07b,graham09}.  The observed projected distances from what has been argued are the plausible hosts, if true, also holds important ramifications for the sort of viable progenitors \\citep{bp06}.  The large offsets seen from early-type hosts would seem to be at odds with progenitor systems with small systematic kicks \\citep[such as in globular clusters][]{Grindlay06,lrr07,rion}, although with such large physical offsets the possibility remains that the association with the putative host is coincidental. On the other hand, based on the small offsets from some low-mass galaxy hosts \\citep{Prochaska05,as06,Berger09,bp06}, SGRB progenitors cannot all have large systematic kicks at birth and inhere large delay times from formation.  Making more quantitative statements about the nature of the progenitor systems is not only hampered by small number statistics but also from the lack of robust predictions of the distribution of merger sites. These distributions depend sensitively on the gravitational potential of the host, which until now has been assumed to be static, and the compact binary formation properties, especially the systematic kick velocity.  Here we have refined such calculations to incorporate the temporal evolution of the host's gravitational potential as well as that of its nearby neighbors self-consistently using cosmological simulations of structure formation. This results in diverse predictions of offsets and compact binary demographics even in the simplest case of a kick velocity distribution \\citep[here assumed to be in excess of $200\\;{\\rm km\\;s^{-1}}$ in order to explain the observed parameters of double neutron star systems;][]{fk97} whose properties do not vary with the initial binary separation.  Two important predictions stand out. First, the merger site distributions computed in this way are more diffusively located with respect to their putative hosts. In a field environment, for example, the distribution of merging sites can extend out to a distance of a few Mpc. This is more severe for those host galaxy halos that were unable to effectively retain most of its own compact binary population at birth. Second, the degree of mixing between neighboring compact binary populations depends on galactic environment. In a cluster, for example, the mixing of the various compact binary populations is severe as a result of the high merger activity and increases dramatically with time. At $z=0$, in a cluster environment, the probability of finding a foreign coalescing compact binary system is equal or higher at all radial distances than finding one originating from the massive central halo. As a result, the closest galaxy at the time of binary coalescence and possibly SGRB occurrence may not to be the one where the compact binary system originated from.  Of course, our basic model is rather simple since we assume a single epoch of star formation and a simple star formation recipe. Also different distributions of kick velocities should be considered. We do not expect that our qualitative results will change dramatically as our first, more general, results indicate.  It is evident form the discussion above that assuming a large (already evolved) host galaxy at the time of compact binary formation thus severely overestimates the binary retention fraction and the concentration of their merging site distribution. This implies that SGRB redshift estimates based solely on the nearest galaxy in projection can be very inaccurate, if progenitor systems inhere large systematic kicks at birth. Interpretations on the nature of the SGRB progenitor using the stellar and mass properties of the nearest galaxy in projection as established by the afterglow location must therefore be regarded with suspicion.  Finally, it should be noted that a direct comparison with model predictions is still impeded by the possibility of an ambient density bias \\citep{Bloom05,Lee05} where SGRBs are more likely to be found in denser gas regions and , as a result, we could be missing a population of bursts with large systematic kicks."
    },
    "0910/0910.2379_arXiv.txt": {
        "abstract": "A method is proposed for constraining the Galactic gravitational potential from  high precision observations of the phase space coordinates of a system of relaxed tracers. The method relies on an ``ergodic'' assumption that the observations are representative of the state of the system at any other time. The observed coordinates serve as  initial conditions  for moving  the tracers forward in time in an assumed model for the gravitational field. The validity of the model is assessed  by the statistical equivalence  between the observations and the distribution of tracers at randomly selected times. The applicability of this ergodic method is not restricted by any assumption on the form or symmetry of the potential.  However, it requires high recision observations as those that will be obtained from missions like SIM and GAIA. ",
        "introduction": "Measurements of velocities of cosmological objects are a classic probe of the mass distribution on all scales. It was the motions of individual galaxies in galaxy clusters which first showed that luminous mater contributed only a small fraction to the total mass in clusters \\cite[]{Zwicky37}, implying the existence of dark matter. The combined mass of the Milky Way (or the Galaxy)  and M31 is constrained from the observed relative  motion between the two galaxies and the requirement that the initial distance between their respective centers of mass vanishes near the Big Bang \\cite[]{KW59}. Line-of-sight velocities of other galaxies in the Local Group  of galaxies also are used to estimate its mass by the condition of vanishing initial distances \\cite[]{Peebles89}. On scales 10s of Mpcs, peculiar motions (deviations of Hubble flow) of galaxies constrain the global mass density in the Universe \\cite[e.g.][]{Nusser08}.    Determining the mass distribution in the Milky Way is particularly important. There is ample information on the  baryonic content  of the Galaxy which could be modeled in detail only if the dark matter distribution is known. The rotation curve of the Galaxy is limited to distances smaller than 20 kpc and does not provide any information on deviations from spherical symmetry of the halo.  The mass distribution at larger distances, motions of Galactic satellite galaxies, globular clusters and stars are invoked \\cite[e.g.][]{sakamoto03}. Constraints on the Galactic mass from these tracers are derived from  the condition that the  observed speeds of Galactic objects do not exceed the escape velocity. This approach yields only a lower limit and is mostly  sensitive to the  highest velocity objects. The vast majority of the sample objects play no role in deriving the mass limit.  Alternatively, one could adopt a Bayesian likelihood formalism in which the phase space distribution function is assumed to follow certain form which could be matched with the observations to probe the Galactic potential field  \\cite[e.g.][]{LT87,Koch96,WE99} .  Proper motions, resulting from  velocities perpendicular to the line-of-sight, are currently measured only for nearby tracers \\cite[e.g.][]{sakamoto03}. Therefore, most current mass estimates rely on the measured line-of-sight motions of individual components.  During the next few years, accurate proper motions for a large sample of Galactic tracers are expected to be measured by the space missions {\\it Global Astrometry Interferometer for Astrophysics } (GAIA) \\cite[]{LP96} and {\\it Space Interferometry Mission } (SIM) \\cite[]{rr1}. Even accurate phase space information require additional assumptions in order to constrain the Gravitational potential \\footnote{The gravitational force field is equal to the acceleration rather than velocities. Measuring the acceleration of a tracer with orbital period $\\td$ over an observing time $\\tobs$ requires astrometry with  angular resolution a factor $\\tobs/\\td$ higher than the precision needed for velocities. Since $\\tobs\\sim$ a few years  while $\\td\\sim$  a few Gyrs, the task is out of reach in the near future. }. We present here a general method which relies on high precision measurements of positions and velocities  of  tracers. The method assumes that the tracers have reached  an equilibrium  state in the Galactic gravitational field. ",
        "conclusions": "\\label{sec:Conclude}  The ergodic method presented here  requires a parametric functional form for the Gravitational potential, but does not impose  any special symmetry on the mass distribution. The method assumes that the observations of a class of tracers are spatially complete. Tracers with observed distances smaller than  $d_0$, could be present beyond $d_0$ in snapshots at later times. Therefore, Observational selection against tracers at distances $>d_0$ will make the statistical comparison between observations and snapshots at later times extremely difficult. The completeness is not too demanding a condition for tracers like globular clusters and Galactic satellites for which future observations should be accurate enough for quite  large Galactic  distances.   We have presented only partial testing of the method, with only a one parameter family for the form of the Galactic potential. Method is able to constrain the parameter to a good accuracy with measurements errors that are even   larger than those expected to be  achieved by future data. The tests show that the method could provide unbiased constraints on the Gravitational potential, but  a more elaborate testing which includes a more realistic treatment of the errors should be done. Observations will  likely assign distance and velocity measurement errors to tracers  on an individual basis. Therefore,  random errors and systematic biases tailored to the specific sample of tracers used by the method could be determined robustly.  The accuracy of the method is mainly limited by the number of tracers. The most obvious tracers  are globular clusters and Galactic satellites. Our Galaxy includes 158 known globular clusters, and 23 known satellites \\cite[e.g.][]{SG07}. Distance measurements of RR Lyrae stars from their period-luminosity relation will be greatly improved by  GAIA and SIM   calibration of the zero point using a nearby sample of these stars. Therefore, luminous halo RR Lyrae stars could significantly enlarge the sample of tracers.  The method requires a   system of tracers in dynamical equilibrium in the current Galactic potential. A pre-requist for dynamical equilibrium is that any recent changes in the Galactic potential  must have occurred on a time scale longer that the dynamical time of the system of tracers. Spectroscopic studies stars in the Galaxy do not present evidence for substantial mergers in the last 8Gyr \\cite[]{GWN,Helmi}, implying a  nearly static gravitational potential.   We have demonstrated that  adiabatic growth of the Galactic disk is not expected to pose a problem for the implementation of the method. The effect of any deviation from that should be modeled.  One issue which is exclusive to using disk stars as tracers is  the effect of transient perturbations on the dynamics of those stars \\cite[e.g.][]{FAM05,QUIL}. However,  the transient effects cause velocity perturbations at the level of ~10 km/s which is much smaller that the total velocities of tracers.  Observational uncertainties are larger than that and do not seem to cause significant biases in the results as indicated by the tests described above. Another issue to be considered is halo substructure  which, in principle, could act as a stochastic component in the gravitational potential. Such a component is extremely difficult to model in the method proposed here. We offer the following argument demonstrating that substructure should not have an important effect on the long term dynamics of tracers. In the limit of fast encounters, a tracer passing a substructure at distance $b$ will change its velocity by $V_1\\approx g(b)b/V$ where $V$ is the relative velocity and $g\\approx G M/b^2$ is the gravitational force field of the substructure assuming that $b$ is larger than its tidal radius. Performing the usual summing in quadratures over encounters occurring in one orbital time we get that the  r.m.s change $<V_1^2>^{1/2}\\approx M\\sqrt{\\bar n R}/V$ where $\\bar n$ is  the number density of substructures and $R$ is the distance travelled by the tracer in one orbital time. Taking $V^2=G M_{_{\\rm MW}}(R)/R$ with $M_{_{\\rm MW}}(R)$  the mass of the Galactic halo within radius $R$, we get the condition $\\sim M_{_{\\rm MW}} \\sqrt{\\bar n R^3}>M_{_{\\rm MW}}$ for having $<V_1^2>^{1/2}\\sim V$. This means that neither single encounters nor collective  stochastic effects  can dominate the  long term evolution of tracers even if the fraction of mass in substructures is large which is contrary to  recent simulations \\cite[e.g.][]{Colombi} which show that the smooth component greatly dominates the mass of Galactic size halos.   When details of future data from SIM and GAIA become available, all robust information about the distribution of baryons in the Galaxy should be used \\cite[e.g.][]{ROBIN03} in order to place tight constraints on the Galactic dark matter. The validity of the method should be tested with mock data that match the observations as much as possible and with the best possible available Galactic models."
    },
    "0910/0910.0183_arXiv.txt": {
        "abstract": "Based on ideas by Woodward et al., a subgrid scale model that is applicable to highly compressible turbulence is presented. Applying the subgrid scale model in large eddy simulations of forced supersonic turbulence, the bottleneck effect is largely reduced and, thereby, approximate scaling laws can be obtained at relatively low numerical resolution. In agreement with previous results from PPM simulations without explicit subgrid scale model, it is found that the energy spectrum function for the velocity field with fractional density-weighing, $\\rho^{1/3}\\vect{u}$, varies substantially with the forcing, at least for the decade of wavenumbers next to the energy-containing range. Consequently, if universal scaling of compressible turbulence exists, it can be found on length scales much smaller than the forcing scale only. ",
        "introduction": "Large eddy simulations (LES) are of great utility to engineers and atmospheric scientists, but not commonly used in computational astrophysics. While terrestrial turbulence is incompressible or weakly compressible in most cases, astrophysicists often deal with highly compressible turbulence. Since the stability of numerical solvers for the equations of compressible gas dynamics requires energy dissipation that is intrinsic to the numerical scheme, gas-dynamical simulations at high Reynolds numbers are considered as \\emph{implicit} large eddy simulations. However, it was shown that numerical schemes such as the piecewise parabolic method (PPM) of \\citet{ColWood84} entail undesired properties of the numerical solutions such as the bottleneck effect, i.~e., an unphysical enhancement of spectral power in the high-wave number range \\citep[e.~g.,][]{KritNor07}. This led \\citet{WoodPort06} to the idea to couple an \\emph{explicit} subgrid-scale (SGS) model to PPM, and they applied the method to decaying transonic turbulence.  In this article, I briefly outline an SGS model that is valid in the supersonic regime. As a first application, I present results from LES of forced supersonic isothermal turbulence with root mean square Mach number between $5$ and $6$, which can be interpreted as an idealized model for cold turbulent gas in the interstellar medium \\citep[see][]{KritNor07,FederDuv09}. Computing turbulence energy spectra from these LES, I find that the SGS model largely reduces the bottleneck effect in comparison to plain PPM simulations. The results shed light on the question of the universality of compressible turbulence, which is important for the theory of star formation in molecular clouds \\citep[see, for example,][]{ElmeSca04}. ",
        "conclusions": "The bottleneck effect, which distorts turbulence energy spectra toward high wavenumbers in PPM simulations without explicit SGS model, is mostly compensated by the SGS turbulence stress in LES. This property is clearly favorable for the determination of scaling laws from the spectra. Moreover, the SGS model for compressible turbulence presented in this article will be valuable in simulations of complex astrophysical flows, where a large range of scales covers energy-injecting physical processes, while only a small portion of the inertial subrange can be resolved. A first example for such an applications has recently been presented by \\citet{MaierIap09}  Based on new data for hydrodynamical as well as magnetohydrodynamical turbulence, \\citet{KritUst09} put forward the hypothesis that the universality of compressible turbulence becomes manifest in the scalings of $\\rho^{1/3}\\vect{u}$. The spectral indices inferred from the $512^{3}$ LES presented in this article, however, do not support this hypothesis. There are two interpretations of these results. One interpretation presumes that the scalings in the range of wavenumbers next to the energy-containing range reflect genuine inertial-range properties of supersonic turbulence. Then the turbulence energy spectra clearly exhibit a dependence of these properties on the forcing \\citep[also see][]{SchmFeder08}. As proposed by \\citet{FederDuv09}, the prominent flattening of the spectrum function $E_{1/3}(k)$ at high wavenumbers that is observed for compressive forcing (see right plot in Fig.~\\ref{fg:spect_512}) might be caused by a transition to the weakly compressible regime at the sonic wavenumber, for which the turbulent velocity fluctuations are comparable to the speed of sound. On the other hand, it is possible that compressible turbulence with universal scaling can only develop on length scales much smaller than forcing scale if the forcing is mostly compressive. In order to settle the question which interpretation is correct, the combined influence of varying the Mach number of the flow and choosing different mixing ratios of solenoidal and compressive forcing components has to be investigated. In this regard, the SGS model presented in this article can help to explore the parameter space."
    },
    "0910/0910.5590_arXiv.txt": {
        "abstract": "{We present new optical long-slit data along 6 position angles of the bulge region of M31. We derive accurate stellar and gas kinematics reaching 5 arcmin from the center, where the disk light contribution is always less than 30\\%, and out to 8 arcmin along the major axis, where the disk makes 55\\% of the total light. We show that the velocity dispersions of McElroy (1983) are severely underestimated (by up to 50 km/s). As a consequence, previous dynamical models have underestimated the stellar mass of M31's bulge by a factor 2. As a further consequence, the light-weighted velocity dispersion of the galaxy grows to 166 km/s and to 170 km/s if also rotation is taken into account, thus reducing the discrepancy between the predicted and measured mass of the black hole at the center of M31 from a factor 3 to a factor 2. The kinematic position angle varies with distance, pointing to triaxiality, but a quantitative conclusion can be reached only after simultaneous proper dynamical modeling of the bulge and disk components is performed. We detect gas counterrotation near the bulge minor axis. We measure eight emission-corrected Lick indices. They are approximately constant on circles.  Using simple stellar population models we derive the age, metallicity and $\\alpha$-element overabundance profiles. Except for the region in the inner arcsecs of the galaxy, the bulge of M31 is uniformly old ($\\ge 12$ Gyr, with many best-fit ages at the model grid limit of 15 Gyr), slightly $\\alpha$-elements overabundant ($[\\alpha/Fe]\\approx 0.2$) and at solar metallicity, in agreement with studies of the resolved stellar components.  The predicted u-g, g-r and r-i Sloan color profiles match reasonably well the dust-corrected observations, within the known limitations of current simple stellar population models.  The stellar populations have approximately radially constant mass-to-light ratios ($M/L_R\\approx 4-4.5M_\\odot/L_\\odot$ for a Kroupa IMF), in agreement with stellar dynamical estimates based on our new velocity dispersions. In the inner arcsecs the luminosity-weighted age drops to 4-8 Gyr, while the metallicity increases to above 3 times the solar value. Starting from 6 arcmin from the center along the major axis, the mean age drops to $\\le 8$ Gyr, with slight supersolar metallicity ($\\approx +0.1$ dex) and $\\alpha-$element overabundance ($\\approx +0.2$ dex), for a mass-to-light ratio $M/L_R\\le 3M_\\odot/L_\\odot$.  Diagnostic diagrams based on the [OIII]/H$\\beta$ and [NI]/H$\\beta$ emission line equivalent widths (EWs) ratios indicate that the gas is ionized by shocks outside 10 arcsec, but an AGN-like ionizing source could be present near the center. We speculate that a gas-rich minor merger happened some 100 Myr ago, causing the observed minor axis gas counterrotation, the recent star formation event, and possibly some nuclear activity.} ",
        "introduction": "\\label{sec_intro} This is the first of two papers presenting new optical spectra for the bulge of M31 to study its stellar populations and assess its triaxiality through dynamical modeling. Here we present the new data and constrain the stellar populations.  In the past 50 years papers studying the dynamics of our neighbour galaxy M31 have been published on a regular basis, discussing gas kinematics, both by optical spectroscopy (Boulesteix et al. \\cite{Boulesteix87}, Pellet \\cite{Pellet76}), and in HI (Kent \\cite{Kent89a}, Braun \\cite{Braun91}, Chemin et al. \\cite{Chemin09} and references therein), stellar kinematics concentrating on the central regions to probe the black hole dynamics (Bender et al. \\cite{Bender05}) or considering the whole bulge (McElroy \\cite{McElroy83}). The data are used to construct dynamical models of the galaxy (Widrow et al. \\cite{Widrow03}, Klypin, Zhao and Somerville \\cite{Klypin02}) and possibly probe the tridimensional distribution of its stellar components.  The question of the triaxiality of M31 bulge has been posed early on (Stark \\cite{Stark77}, Gerhard \\cite{Gerhard86}) and is of significant importance for the understanding of M31, but a definitive quantitative modeling of both photometry and kinematics is still missing. A bar could also be present (Athanassoula \\& Beaton \\cite{Athanassoula06}, Beaton et al. \\cite{Beaton07}).  Moreover, investigations of the stellar populations of the central regions of M31 through the measurement of Lick indices have been performed (Davidge \\cite{Davidge97}). They indicate the presence of a young and metal rich population in the inner arcsecs of the galaxy.  Parallely, studies of the resolved stellar population of the bulge of M31 have assessed that the global stellar population of the M31 bulge must be as old as the bulge of Milky Way, resolving previous claims of younger ages as due to crowding problems (Stephens et al.  \\cite{Stephens03}).  Two considerations convinced us of the necessity to collect new optical spectroscopic information for the bulge of M31, in addition to the old age of the dataset of McElroy (\\cite{McElroy83}). The first one is the start of PAndromeda, an extensive monitoring campaign of M31 with the PanSTARRS-1 telescope and camera system (Kaiser \\cite{Kaiser04}), that is in principle able to deliver hundreds of pixel lensing events, probing both bulge and disk regions. Detailed stellar population and dynamical models, based on accurate spectral information, are needed to interpret these events as due to a compact baryonic dark matter component (the so-called MACHOs) rather than self-lensing of stellar populations (Kerins et al. \\cite{Kerins01}, Riffeser et al.  \\cite{Riffeser06}).  The second is the development of new modeling techniques of both simple stellar populations and stellar dynamical systems. On the one hand, the interpretation of Lick indices (Worthey et al \\cite{Worthey94}) in terms of the most recent simple stellar population models (Maraston \\cite{Maraston98}, \\cite{Maraston05}) that take into account the variation of $[\\alpha/$Fe] (Thomas, Maraston and Bender \\cite{TMB03}), allows the accurate determination of the stellar population ages, metallicities and overabundances, and therefore the prediction of stellar mass-to-light ratios. On the other hand, new dynamical modeling codes, like N-MAGIC (De Lorenzi et al.  \\cite{DeLorenzi07}) allow the flexible dynamical modeling of triaxial structures, optimally exploiting the information contained in the line-of-sight velocity distributions that modern programs for the analysis of the galaxy optical spectra are able to extract (Bender, Saglia and Gerhard \\cite{BSG94}), well beyond the mean velocities and velocity dispersions of McElroy (\\cite{McElroy83}).  In the following we discuss our new spectroscopic observations of the bulge of M31. A future paper (Morganti et al. \\cite{Morganti09}) will report on the dynamical modeling. In Sect. \\ref{sec_obs} we present the observations and the data reduction.  In Sec. \\ref{sec_kinematics} we derive the stellar and gas kinematics and the strengths of the absorption and emission lines. In Sec. \\ref{sec_mod} we discuss the modeling of the new spectroscopic data. We analyze the stellar population in Sect. \\ref{sec_stelpop} and discuss previous axisymmetric dynamical models of the bulge of M31 in Sect.  \\ref{sec_dyn}. Sect. \\ref{sec_ionize} considers the possible excitation sources compatible with the observed emission line EW ratios. We draw our conclusions in Sect.  \\ref{sec_conc}. ",
        "conclusions": "\\label{sec_conc}  We have presented new optical spectroscopic observations of the bulge of M31. We have measured the stellar and gas kinematics, emission line strength ratios and absorption Lick indices profiles along 6 position angles out to distances of 5 arcmin from the center. Along the major axis we probed regions out to 8 arcmin. We show that the old kinematics of McElroy (\\cite{McElroy83}) provides too small velocity dispersions (by up to 30\\%), therefore biasing the dynamical modeling towards too small bulge stellar masses (by a factor 2).  Moreover, the new higher averaged velocity dispersion predicts a mass for the central supermassive black hole of M31 that is only a factor 2 below what measured. The new velocity dispersion profiles are now in better agreement with axisymmetric dynamical models with large bulge mass-to-light ratio (Widrow et al. \\cite{Widrow03}), that now match the values derived from stellar population models ($M/L_R\\approx 4-4.5 M_\\odot/L_\\odot$, see below). This implies an upward revision of the predicted self-lensing microlensing event rate of Widrow et al. (\\cite{Widrow03}), and Riffeser et al.  (\\cite{Riffeser06}), that are based on lower stellar mass-to-light ratios.  The inner ($r\\le 100$ arcsec) bulge is slowly rotating, with a $V/\\sigma\\approx 0.2$. At distances from the center larger than $\\approx 100$ arcsec the measured kinematics becomes increasingly influenced by the rapidly rotating stellar disk. Therefore, the observed variation of the kinematic position angle is suggestive of bulge triaxiality, but needs a proper dynamical modeling of both disk and bulge components to be quantified. The measured gas kinematics confirms the well studied large scale disk rotation. However, a more complex structure, with gas minor axis counter-rotation, is detected in the inner bulge. This might be evidence for a (recent) minor merger, possibly connected to the younger stellar population detected in the inner  arcsecs of the galaxy discussed below.  The analysis of eight Lick index profiles shows that the bulge of M31 is old, of solar-metallicity and a factor 2 overabundant in $\\alpha$-elements, in agreement with studies of its resolved stellar component (Stephens et al. \\cite{Stephens03}). The line indices and stellar population parameters appear approximately constant on circles, i.e. their isocontours are rounder than the galaxy isophotes, as seen in many ellipticals and bulges (Kuntschner et al. \\cite{Kuntschner06}; Falc\\'on-Barroso et al.  \\cite{Falcon06}). Together with the derived old ages, this confirms that the stellar disk out to 5 arcmin from the center is either old (similar to what found for other spiral galaxies, Peletier and Balcells \\cite{Peletier96}) or not sufficiently probed by the spectral features considered here. However, we do detect smaller ages ($\\le 8$ Gyr) along the major axis at distances $\\ge 6$ arcmin. The u-g SLOAN colors predicted from our stellar population analysis match reasonably well with the observed ones. The redder colors g-r and r-i are systematically offset, a well known problem of the flux calibration of current simple stellar population models (Maraston et al. \\cite{Maraston09}). The derived mass-to-light ratios (in the Johnson R band and with a Kroupa IMF we get $M/L\\approx 4-4.5$) agree with the dynamical estimates (see above). They drop to $\\le 3 M_\\odot/L_\\odot$ along the major axis at distances $\\ge 6$ arcmin, where the disk light starts to dominate.  In the inner arcsecs the situation changes and a population with a light-weighted younger age ($\\approx 8$ Gyr inside the seeing disk or 2 arcsec, with values as low as 4 Gyr) and metal richer ($\\approx 3$ times solar) appears.  This agrees with the findings of Davidge (\\cite{Davidge97}). In addition, the emission line EW ratios [OIII]/H$\\beta$ increase in this region. Their values are compatible with being excited by shocks in the main body of the bulge, but near the center they increase to levels pointing to the presence of an AGN-like photo-ionizing source. Combined with the detection of counterrotating gas along the minor axis of the galaxy (see above), this suggests that a gas-rich minor merger probably happened some 100 Myr ago, that triggered an episode of star formation and possibly boosted the nuclear activity of the central supermassive black hole of M31.  We estimate how broad a range of star burst ages and masses can be to give the measured mean value of 8 Gyr in the inner 2 arcsec, when superimposed to the old bulge stars background. To this purpose we compute composite spectra of an old (12.6 Gyr) plus a young (from 100 Myr to 4 Gyr) simple stellar population, using the Vazdekis (\\cite{Vazdekis99}) library, and measure their SSP ages through the analysis of the Lick indices observed here. We find that global ages smaller than 8 Gyr are found when considering a young component younger than $\\approx 600$ Myr and a mass fraction lower than 10\\%. Higher mass fractions are possible for older ages. From Kormendy and Bender (\\cite{Kormendy99}), inside an aperture of 2 arcsec diameter we measure a V mag of 12.5, or $5\\times 10^6 L_\\odot$. In this region we estimate $M/L_V\\approx 4 M_\\odot/L_\\odot$ and therefore an enclosed mass of $2\\times 10^7 M_\\odot$. As a consequence, $\\approx 10^6 M_\\odot$ of some 100 Myr old stars would be needed. Note that in the inner nucleus of M31, at fractions of an arcsec, a disk of 200 Myr old stars is found (Bender et al.  \\cite{Bender05}), with a mass of $\\approx 4200 M_\\odot$. Of course the spatial resolution of our spectra is too low to probe this scale.  Moreover, our result, based on the H$\\beta$ line as an age tracer, depends heavily on the details of the emission correction and might be affected by the bad column's interpolation (see Sect. \\ref{sec_obs}). Spectra of the higher order Balmer lines, taken with better seeing, are needed to improve our conclusions.  In a second paper (Morganti et al. \\cite{Morganti09}), a dynamical model of the data presented here, that takes into account the contributions of the bulge and disk components, will assess in a quantitative way the bulge triaxiality issue and will give a proper estimate of the microlensing event rates due to self-lensing."
    },
    "0910/0910.2715_arXiv.txt": {
        "abstract": "We present the first results of a program to characterize the disk and envelope structure of typical Class~0 protostars in nearby low-mass star forming regions. We use \\textit{Spitzer} IRS mid-infrared spectra, high resolution CARMA 230~GHz continuum imaging, and 2-D radiative transfer models to constrain the envelope structure, as well as the size and mass of the circum-protostellar disk in Serpens FIRS~1. The primary envelope parameters (centrifugal radius, outer radius, outflow opening angle, and inclination) are well constrained by the spectral energy distribution (SED), including \\textit{Spitzer} IRAC and MIPS photometry, IRS spectra, and 1.1~mm Bolocam photometry.  These together with the excellent $uv$-coverage ($4.5-500$~k$\\lambda$) of multiple antenna configurations with CARMA allow for a robust separation of the envelope and a resolved disk. The SED of Serpens FIRS~1 is best fit by an envelope with the density profile of a rotating, collapsing spheroid with an inner (centrifugal) radius of approximately 600~AU, and the millimeter data by a large resolved disk with $\\mdisk\\sim1.0~ \\msun$ and $\\rdisk\\sim 300$~AU. These results suggest that large, massive disks can be present early in the main accretion phase.  Results for the larger, unbiased sample of Class~0 sources in the Perseus, Serpens, and Ophiuchus molecular clouds are needed to determine if relatively massive disks are typical in the Class~0 stage. ",
        "introduction": "Protostars build up their mass by accreting material from a dense protostellar envelope, presumably via a rotationally supported circum-protostellar accretion disk. Disk formation is a natural result of collapse in a rotating core, but it is not known how soon after protostellar formation the disk appears, or how massive it is at early times. Theory suggests that centrifugally supported disks should start out small (radius less than $10$~AU), and thus low mass, and grow with time \\citep*{tsc84}. Unstable or magnetically supported disks, however, could be much larger (radii up to 1000~AU in the magnetically supported case; \\citealt{gs93}), and thus more massive at early times.  The remnants of these protostellar accretion disks are easily observed in more evolved phases (e.g. TTauri stars), but given that the majority of mass is accreted during earlier embedded phases, understanding disks at early times is critical.  Directly observing disks during this main accretion phase is quite difficult, however, as they are hidden within dense, extincting protostellar envelopes. The structure of the envelope at small radii is another important characteristic of main accretion phase protostars that is similarly difficult to directly observe. Disk growth or the presence of a binary companion may clear out the inner region of the envelope early on, as inferred for the binary Class~0 source IRAS 16293-2422 by \\citep{jorg05}.  There has been a recent push toward detecting disks in more embedded objects, with many now known and roughly characterized in Class I protostars \\citep[e.g.][]{loon00,jorg05b,eisner05,aw07}, and a few detected in the earlier Class 0 stage \\citep[e.g.][]{chan95,harv03,brown00,loon00}. \\footnote{We use definitions of Class 0, Class I, and Class II \\citep{andre94} based on the bolometric temperature \\citep{ml93,chen95}: $\\tbol \\le 70$~K (Class~0); 70 K$< \\tbol \\le 650$~K (Class~I); 650 K$<\\tbol \\le 2800$ K (Class~II).  We further assume that classes correspond to an evolutionary sequence \\citep[e.g.][]{rob06}: in Class~0 the protostar has accreted less than half its final mass ($M_{*} < \\menv$), in Class~I $M_* > \\menv$, and in Class~II the envelope has dispersed, leaving only a circum-stellar disk.} Most previous detailed studies have been limited to the most well-known or brightest Class 0 sources, however, due to instrumental limitations and a lack of large unbiased target samples.  The ongoing Submillimeter Array survey of low-mass protostars \\citep{jorg07,jorg09} is a notable exception.  With recent large surveys of nearby molecular clouds at mid-infrared and (sub)millimeter wavelengths it is now possible to define complete samples of Class 0 protostars based on luminosity or envelope mass limits \\citep[e.g.][]{hatch07,jorg08,dun08,enoch09,evans09}.  \\begin{figure*}[!ht] \\vspace{-0.2in} \\begin{center} \\includegraphics[width=4.8in]{f1.eps} \\vspace{-0.2in} \\caption {\\textit{Spitzer} $24~\\micron$ image of the immediate environment of Serpens FIRS~1, in the Serpens main core.  Bolocam 1.1~mm continuum contours are overlaid. Submillimeter source designations for the brightest \\citet{casali93} sources are indicated.  The nearest embedded protostar to FIRS~1 is approximately $45\\arcsec$ or $11000$~AU away (Ser-emb 12; \\citealt{enoch09}), and the nearest YSO is $25\\arcsec$ or 6000~AU away (c2d 142; \\citealt{harv07}). \\label{genfig}} \\end{center} \\end{figure*}  We have recently begun a campaign to characterize disk properties in a large, uniform sample of Class 0 protostars in nearby low mass star-forming regions (M. Enoch et al. 2009, in preparation). Our study is based on the complete (to envelope masses $\\gtrsim0.1~\\msun$) sample of 39 Class 0 protostars in the Serpens, Perseus, and Ophiuchus molecular clouds, identified by \\citet{enoch09} by comparing large-scale \\textit{Spitzer} IRAC and MIPS and Bolocam 1.1~mm continuum surveys of the three clouds.  Combining Spitzer IRS mid-infrared (MIR) spectra and high resolution CARMA 230~GHz continuum imaging with radiative transfer modeling of this sample will help to address several fundamental questions about the structure and evolution of the youngest protostars: 1) How soon after the initial collapse of the parent core does a circum-protostellar disk form? 2) What fraction of the total circum-protostellar mass resides in the disk, and does this fraction vary with time?  3) Are large ``holes'' in the inner envelope, such as that found for IRAS~16293 by \\citet{jorg05}, common at early times?   MIR spectra and millimeter maps provide complementary approaches to these questions. The amount of flux escaping at $\\lambda \\lesssim 50~\\micron$ from deeply embedded sources is very sensitive to the opacity close to the protostar, and thus the envelope structure \\citep[e.g.][]{jorg05}. While the MIR flux is insensitive to disk properties, high resolution millimeter continuum mapping can directly detect emission from dust grains in the disk.  Millimeter observations with excellent $uv$-coverage, combined with radiative transfer models, can separate the disk from the envelope and constrain the disk mass and size.  Our ultimate goal is to characterize the disk mass, size, and inner envelope structure of typical low-mass Class~0 protostars, and to quantify any trends with evolutionary indicators. In this initial paper we present results for Serpens FIRS~1, a well known Class 0 source, which will serve as a test case for the full program. ",
        "conclusions": ""
    },
    "0910/0910.0350.txt": {
        "abstract": "{ An overview of some recent developments in inhomogeneous models is presented.  As the volume and precision of cosmological data improves, it will become more and more essential to understand the non-linear behaviour of the Einstein field equations.  This requires the study of exact inhomogeneous solutions, including their density distributions, their evolution, their geometry, and their causal structure.  Observations are strongly affected by the detailed geometry and evolution of a model, and therefore interpretation of observations depends on understanding them.  It is generally assumed the universe is homogeneous if averaged over large enough scales, but to actually prove this is so, will require the assumption to be relaxed, and a rigorous inhomogeneous approach to be applied.  Though the \\LT\\ metric has long been used for models of spherical inhomogeneities, there have been a number of new results, including a variety of methods for creating models with specific properties, and their application to cosmic structures on several different scales.  Interest in the Szekeres metrics is on the increase, and the quasi-spherical metric was recently used to model specific cosmic structures for the first time.  The quasi-planar and quasi-hyperspherical metrics have been hardly studied until recent work invesigated their physical and geometric properties.  There is enormous scope for work with these metrics. }  \\FullConference{5th International School on Field Theory and Gravitation,\\\\ April 20-24, 2009\\\\ Cuiab\\'{a} city, Brazil}  \\tableofcontents  \\begin{document} ",
        "introduction": "Why study inhomogeneous models?  The real universe is very lumpy.  To properly understand what we see, we should apply all possible methods: perturbation theory, $N$-body newtonian simulations, and exact inhomogeneous metrics --- each has its domain of validity.  Inhomogeneous metrics have the advantage that they are fully non-linear and relativistic solutions of the Einstein field equations (EFEs).  The assumption of homogeneity has become so well established, that it has become all-pervasive.  But now, with so much data coming in, it's time to test homogeneity.  Cosmological data reduction relies heavily on the Robertson-Walker (RW) metric --- we need to beware of a circular argument.  It will be a significant challenge to check which of our well-known results actually depend on the assumption of homogeneity, and to re-derive them all without that assumption.  Here I will present a selection of results in inhomogeneous cosmology, especially work done with Lu, McClure, Krasi\\'{n}ski, Bolejko, C\\'{e}l\\'{e}rier, Alfedeel, Mustapha, Ellis and others, but I won't try to be comprehensive.  I'll attempt to provide the basics, and thereby promote the use of inhomogeneous metrics for the study of cosmological problems.  Inhomogeneous metrics will become more important as the amount and accuracy of cosmological data increases, and more precise analysis is needed, so there are plenty of opportunities for good research. ",
        "conclusions": "The universe is of course very inhomogeneous on many scales.  To fully understand how these structures evolve, and properly analyse our observations will require the non-linearity of exact inhomogeneous metrics.  Up to now, homogeneity has been assumed, and was key to making progess.  In the age of precision cosmology, we should thoroughly test this assumption and quantify how good an approximation it is on each scale.  Nearly all data analysis assumes the RW metric.  To be sure we avoid circular arguments, there is an urgent need to re-do all calculations in a general non-homogeneous metric.  The methods of inhomogeneous cosmology will be an essential component of this endeavour.  \\LT\\ models have produced a wide variety of interesting results, and the investigations are far from exhausted.  The Szekeres models have a lot of flexibility, and can be used to model quite complex structures --- but they have been very little investigated.  There are plenty of opportunities for good research.  \\appendix"
    },
    "0910/0910.1716_arXiv.txt": {
        "abstract": "{The galaxy cluster XMMU~J2235.3$-$2557 (hereafter XMM2235), spectroscopically confirmed at $z=1.39$, is one of the most distant X-ray selected galaxy clusters. It has been at the center of a multi-wavelength observing campaign with ground and space facilities.} {We characterize the galaxy populations of passive members, the thermodynamical properties and metal abundance of the hot gas, and the total mass of the system using imaging data with HST/ACS ($i_{775}$ and $z_{850}$ bands) and VLT/ISAAC (J and K$_S$ bands), extensive spectroscopic data obtained with VLT/FORS2, and deep (196\\,ks) \\Chandra observations. } {\\Chandra data allow temperature and metallicity to be measured with good accuracy and the X-ray surface brightness profile to be traced out to 1\\arcmin\\ (or 500 kpc), thus allowing the mass to be reliably estimated.  Out of a total sample of 34 spectroscopically confirmed cluster members, we selected 16 passive galaxies (without detectable [OII]) within the central 2\\arcmin\\ (or 1 Mpc) with ACS coverage, and inferred star formation histories for subsamples of galaxies inside and outside the core by modeling their spectro-photometric data with spectral synthesis models.} {\\Chandra data show a regular elongated morphology, closely resembling the distribution of core galaxies, with a significant cool core. We measure a global X-ray temperature of $kT = 8.6_{-1.2}^{+1.3}$ keV (68\\% confidence), which we find to be robust against several systematics involved in the X-ray spectral analysis. By detecting the rest frame 6.7 keV Iron K line in the \\Chandra spectrum, we measure a metallicity $Z= 0.26^{+0.20}_{-0.16} \\, Z_\\odot$.  In the likely hypothesis of hydrostatic equilibrium, we obtain a total mass of $M_{\\rm tot}(<1\\,{\\rm Mpc})= (5.9\\pm 1.3)\\times 10^{14}\\,M_\\odot$.  By modeling both the composite spectral energy distributions and spectra of the passive galaxies in and outside the core, we find a strong mean age radial gradient.  Core galaxies, with stellar masses in excess of $10^{11} M_\\odot$, appear to have formed at an earlier epoch with a relatively short star formation phase ($z=5-6$), whereas passive galaxies outside the core show spectral signatures suggesting a prolonged star formation phase to redshifts as low as $z\\approx 2$. } {Overall, our analysis implies that XMM2235 is the hottest and most massive bona--fide cluster discovered to date at $z>1$, with a baryonic content, both its galaxy population and intracluster gas, in a significantly advanced evolutionary stage at 1/3 of the current age of the Universe. } ",
        "introduction": "Over the past two decades, considerable effort has been devoted to discovering ever more distant galaxy clusters using different observational methods \\citep[e.g.][for a review]{2002ARA&A..40..539R}. These studies have been traditionally motivated by cosmological applications of the cluster abundance at high redshift \\citep[e.g.][for a review]{2005RvMP...77..207V}, and also by the use of clusters as laboratories to investigate galaxy evolution. Clusters provide a convenient and efficient way of studying large populations of early-type galaxies, which provide stringent tests on galaxy evolution models in the current hierarchical formation paradigm \\citep{2006ARA&A..44..141R}, because they are the most massive galaxies with the oldest stellar populations (at least out to $z\\sim\\!  2$). Clearly, the higher the redshift the stronger the leverage on theoretical models. In addition, galaxy properties in clusters can be contrasted with those in field surveys, which have multiplied in recent years, thus extending the baseline over which environmental effects can be studied.  X-ray selection of clusters has been central in these studies, as it naturally provides gravitationally bound systems (as opposed to simple overdensities of galaxies) with a relatively simple selection function. Using ROSAT serendipitous surveys, supported by near IR imaging and spectroscopy with 8-10m class telescopes, the redshift envelope was pushed to $z=1.3$, with only 5 clusters discovered at $z>1$, approximately one per square degree \\citep{2002ARA&A..40..539R}.  The extension of the same technique to \\XMM serendipitous surveys has led to the discovery of two clusters at $z>1.3$ to date, XMMU J2235.3$-$2557 at $z=1.39$ \\citep{2005ApJ...623L..85M} (hereafter M05) and XMMXCS~J2215.9$-$1738 at $z=1.46$ \\citep{2006ApJ...646L..13S}. See also \\citet{2008A&A...487L..33L} for a newly discovered, high X-ray luminosity massive cluster at $z\\sim\\! 1$.  With the advent of the {\\it Spitzer} observatory, an alternative and efficient way to unveil distant clusters over large areas (as red galaxy overdensities in the IRAC and optical bands) has been developed. This technique has given notable results, with three clusters spectroscopically confirmed at $z>1.3$ in the IRAC Shallow Survey \\citep[ISCS;][]{2008ApJ...684..905E} and one in the SpARCS survey \\citep{2009ApJ...698.1943W}.  Detailed investigations of galaxy populations in the X-ray luminous clusters at $z>1$ have been conducted with HST/ACS in combination with the VLT and the Keck telescopes. The study of clusters at $z=1.10,\\, 1.24,\\, 1.26,\\, 1.27$ \\citep{2006ApJ...644..759M, 2007ApJ...663..164D, 2009ApJ...690...42M} and the aforementioned systems at $z=1.39$ \\citep{2008A&A...489..981L} and $z=1.46$ \\citep{2009ApJ...697..436H}, have revealed tight red sequences for the early-type galaxies, with scatters only marginally larger than those in local clusters, implying that most of their stellar mass was assembled at $z>3$ with passive evolution thereafter.  While a change in the morphology-density relation of early type galaxies (E+S0 galaxies) has been observed at $z\\sim\\! 1$, elliptical galaxies still dominate the cluster galaxy population up to $z\\sim\\! 1.2$ \\citep{2005ApJ...623..721P,2007ApJ...670..190H}.  A comparison of cluster and field early-types of similar stellar mass has revealed a mild but significant difference between the star formation histories in the two environments \\citep[e.g.][]{2008A&A...488..853G,2008arXiv0806.4604R,2007ApJ...655...30V}, a result which is predicted by current hierarchical galaxy formation models \\cite[e.g.][]{2008ApJ...685..863M}. To date, such comparative studies can only be carried out at $z\\lesssim 1.4$, whereas the existence of a substantial population of old, massive, passively evolving early-type galaxies in the field is now well established up to $z\\sim 2$ \\citep[][]{2006ApJ...649L..71K,2008A&A...482...21C}.  By pushing cluster studies to higher redshifts, where evolutionary time scales become comparable to the age of the Universe at these redshifts, one would expect to detect significant evolutionary effects.  However, as discussed in this paper, this has not been the case so far, even after probing two-thirds of the look-back time, not only for the galaxy populations but also for the thermodynamical properties and chemical enrichment of the hot gas measured with follow-up \\Chandra observations in $z>1$ clusters.  Our current understanding is that relations, such as the red sequence in the color-magnitude diagram, the morphology-density relation, the $L_X-T_X$ relation for the intracluster gas, emerge at $z\\lesssim 2$. For example, the study of a proto-cluster at $z=2.16$ identified around a powerful radio galaxy provided evidence of a forming red sequence \\citep{2008ApJ...680..224Z,2007MNRAS.377.1717K}, which likely takes 1-2 Gyrs to form \\citep{2008A&A...488..853G}.  It is unfortunate that such a transition in the assembly process of galaxy clusters seems to occur in a redshift range where spectroscopic observation are particularly difficult. In this spirit, we have carried out a multi-wavelength study of one the most distant clusters known, XMMU J2235.3$-$2557 (hereafter XMM2235) at $z=1,39$, which was the first distant cluster confirmed (M05) as part of the on-going \\XMM Distant Cluster Project \\citep[XDCP,][]{2005Msngr.120...33B,2008arXiv0806.0861F}.  In this paper, we use spectro-photometric observations of XMM2235 in the optical/near-IR and X-ray bands to characterize its galaxy population (particularly the passive spectroscopic members) and the thermodynamic status of the hot intracluster gas and to measure its total mass. In section \\ref{sec:VLTHST}, we present the VLT and HST data used in this paper, including an extensive spectroscopic campaign which yielded 34 confirmed cluster members. In section \\ref{sec:SFH}, we model the underlying stellar populations of passive galaxies and constrain their star formation histories. In section \\ref{sec:Xray}, we present deep \\Chandra observations of XMM2235, the methods of analysis and the resulting measurements of its temperature, metallicity and mass. In section \\ref{sec:results}, we discuss the results.  $H_0=70\\ {\\rm km}\\ {\\rm s}^{-1}\\,{\\rm Mpc}^{-1}, \\,\\Omega_m=0.3, \\, \\Omega_\\Lambda=0.7$ are adopted throughout this paper. In this cosmology, 1\\arcmin\\ on the sky corresponds to 0.5 Mpc at $z=1.39$ ",
        "conclusions": "\\label{sec:results}  We presented a combination of spectro-photometric data from VLT and HST, as well as deep \\Chandra observations of the X-ray selected cluster XMMU J2235.3$-$2557 at $z=1.39$ and used them to characterize the galaxy populations of passive members, the thermodynamical properties of the hot gas and the total mass of the system.  Despite surface brightness dimming, due to the high redshift of the cluster, the \\Chandra data show extended X-ray emission out to $\\sim\\! 0.5$ Mpc. The X-ray emission has a regular morphology and is clearly elongated in the same way as the distribution of red passive galaxies in the core as well as the major axis of the BCG. An excess of emission in the inner 50 kpc is clearly detected and naturally interpreted as a cool core. The spectral analysis of the \\Chandra data reveals that XMM2235 has a global temperature of $kT = 8.6_{-1.2}^{+1.3}$ keV (68\\% confidence), which we find robust against several systematics involved in the X-ray spectral analysis. If we assume hydrostatic equilibrium, a well justified condition given the relaxed appearance and the canonical $L_X/T_X$ ratio of the cluster, and an isothermal gas distribution, the X-ray surface brightness profile yields a total mass at large radii ($r\\gg 100$ kpc) of $M_{\\rm tot}(<r)\\approx 5.9\\times 10^{14}\\, r({\\rm Mpc})\\, M_\\odot$.  Overall, our analysis implies that XMM2235 is the hottest and most massive bona--fide cluster discovered to date at $z>1$. These findings are corroborated by the weak lensing analysis from deep HST/ACS data \\citep{2009ApJ...704..672J} which provides a mass profile in very good agreement with the X-ray measurement.  The presence of a significant cool core is additional evidence of the advanced dynamical state of the cluster \\citep[e.g.][]{1994ARA&A..32..277F}.  In characterizing the passive galaxy population of XMM2235, we have extended the analysis of the ACS Intermediate Redshift Cluster Survey at $0.8<z<1.3$ \\citep[][]{2009ApJ...690...42M, 2007ApJ...663..164D} to higher redshifts.  Using a large sample of spectroscopically identified cluster members, we inferred the SFHs of passive galaxies in the core ($R_{CL}<100$ kpc) and in the outskirts by modeling their spectro-photometric data with spectral synthesis models. We find a clear contrast in the underlying stellar populations of the two samples. Core galaxies, all with photometric stellar masses in excess of $10^{11} M_\\odot$, appear to have formed at an earlier epoch with a relatively short star formation phase ($z_F=6-4$), whereas passive galaxies outside the core show spectral signatures that suggest a star formation phase that is prolonged to later epochs, to redshifts as low as $z\\approx 2$. The latter are presumably the infalling population of galaxies that are in the process of populating the red sequence.  These results are also consistent with their mean color and scatter of the $J-K_S$ red sequence \\citep{2008A&A...489..981L}.  The on-going star formation is confined to the outskirts of XMM2235, as galaxy members with a detectable [OII] emission line avoid the inner 250 kpc region. Also in this respect, XMM2235 has already reached an advanced evolutionary stage, as star formation is suppressed in galaxies well before they reach the center, similarly to low redshift massive clusters.  We note that the environmental age gradient we find in XMM2235, as well as in other $z\\gtrsim 1$ clusters \\citep[e.g.][]{2009ApJ...690...42M}, is significantly steeper than the relatively mild difference found between {\\sl field} and {\\sl cluster} early-type galaxies of similar stellar masses \\citep{2008A&A...488..853G}. This shows the importance of environmental effects in driving galaxy evolution, in keeping with expectations from hierarchical galaxy formation models \\citep[e.g.][]{2008ApJ...685..863M}.  It would be interesting to study whether the radial age gradient of galaxy populations becomes larger with increasing redshifts. This would require homogeneous galaxy selection criteria in large samples of distant clusters with homogeneous photometric measurements, a task which might be challenging with current data.  At increasing redshifts, the observational effects of different evolutionary histories of galaxies from cluster to cluster should become apparent. Mild differences have been detected so far in the most distant X-ray luminous clusters \\citep[e.g.][in the mean color of the color sequence]{2009ApJ...690...42M} of otherwise very homogeneous populations.  \\cite{2009ApJ...697..436H} carried out a detailed photometric and morphological study of XMMXCS J2215.9-1738 at $z=1.46$, currently the most distant spectroscopically confirmed X-ray luminous cluster. Based on its X-ray temperature and luminosity, as well its velocity dispersion, this cluster is less massive than XMM2235 and appears to be in a younger dynamical state based on its velocity distribution \\citep{2007ApJ...670.1000H}. Another striking difference between the two clusters is the lack of a dominant BCG in XMM2215. However, as in XMM2235, the galaxy population of XMM2215 is dominated by early-type galaxies with ages of $\\sim\\! 3$ Gyr, based on the scatter of its red sequence, implying a major episode of star formation at epochs $z_F\\approx 3-5$.  Unfortunately, it is difficult to directly compare these results with those we obtained for XMM2235 as the mean ages of the early-type galaxies were estimated with different methods. A similar comparative study with optically selected clusters would be very interesting but is still lacking. In general, the study of cluster-to-cluster variance in distant clusters, leading to different evolutionary histories of cluster galaxies, could provide further tests for galaxy evolution models. A robust comparison among different clusters however would require a homogeneous observing campaign with the same set of instruments and filters so as to minimize systematics in photometric measurements and model dependent K-corrections.  Another notable result of our study is the measured metal abundance $Z = 0.26^{+0.20}_{-0.16} \\, Z_\\odot$ of the intracluster medium (ICM) from the detection of the Fe-K line in the \\Chandra spectrum, essentially coming from the cluster core ($R_{CL}\\lesssim 200$ kpc where most of the \\Chandra counts are contained). This extends to even larger look-back times the evidence that the ICM was already significantly enriched in distant clusters, as found in a systematic study of clusters at $z<1.3$ \\citep[][]{2007A&A...462..429B, 2008ApJS..174..117M}, and therefore that metal enrichment of the ICM was mostly complete at early epochs.  In light of the derived SFH of the core red galaxies this is not unexpected. Chemical evolution models of elliptical galaxies \\citep[e.g.][]{2004MNRAS.347..968P} predict that chemical enrichment of massive elliptical galaxies, which experience a short burst of star formation at early epochs, occurs on a time scale of 1 Gyr. Galactic winds distribute metals in the cluster core region on a crossing time scale ($\\lesssim 1$ Gyr). As a result, one would expect the metal enrichment of the ICM to be completed by $z\\gtrsim 2.5$ (2 Gyr earlier than the cluster look-back time of 9 Gyr), consistently with the derived SFH of the core galaxies.  It remains somewhat surprising that such a massive cluster, which was discovered from a relatively small survey area (11 deg$^2$), is found with a baryonic content in such an advanced evolutionary stage at an epoch corresponding to 1/3 of the current age of the Universe, both in terms of its galaxy population and of the hot intracluster gas. Further results on the galaxy populations of XMM2235 as well as its mass distribution, taking advantage of an augmented multi-wavelength data set, will be presented in forthcoming papers.  Given its large mass, XMM2235 will also likely be an ideal target for a new array of Sunayev-Zeldovich (SZ) experiments which are already operational \\citep[e.g.][]{2009ApJ...701...32S} or are becoming available. These observations will add independent, important information on the physics of the ICM and the total mass of XMM2235 and will be useful to calibrate the expectations of a new era of SZ cluster surveys."
    },
    "0910/0910.1466_arXiv.txt": {
        "abstract": "The evolution of AGN in groups and clusters provides important information about how their black holes grow, the extent to which galaxies and black holes coevolve in dense environments, and has implications for feedback in the local universe and at the epoch of cluster assembly. I describe new observations and analysis that demonstrates that the AGN fraction in clusters increases by a factor of eight from the local universe to $z \\sim 1$ and that this evolution is consistent with the evolution of star-forming galaxies in clusters. The cluster AGN fraction remains approximately an order of magnitude below the field AGN fraction over this entire range, while a preliminary analysis of groups indicates that they too undergo substantial evolution. ",
        "introduction": "Many studies over the past decade have presented strong evidence for the coevolution of black holes and galaxies based on samples dominated by the low-density field \\citep[e.g.][and references therein]{hopkins06a,silverman08b}. It is interesting to determine if similar coevolution between AGN and galaxies is present in groups and clusters because the physical processes that drive galaxy evolution, such as the available cold gas to fuel star formation and black hole growth, are substantially different from the field. In addition, AGN in groups and clusters at low-redshift appear to play the critical role in maintaining the temperature of their hot atmospheres \\citep[e.g.][]{mcnamara07}, while AGN heating at the epoch of cluster assembly remains a viable explanation for the minimum entropy level in the intracluster medium.  These questions have motivated my collaborators and I to systematically search for AGN in groups and clusters of galaxies and led to the discovery of large numbers of AGN in these dense environments in the local universe \\cite{martini06,martini07}. In this contribution I summarize some recently published work on the evolution of AGN in clusters of galaxies \\cite{martini09} and present a new, preliminary analysis of the evolution of AGN in the lower-density group environment. I conclude with a brief discussion of several future directions. ",
        "conclusions": "The observations summarized here have shown that the AGN fraction in clusters increases by approximately an order of magnitude from the present to $z \\sim 1$. This evolution is also quantitatively similar to the evolution of the star-forming galaxy fraction in clusters, which suggests that cluster galaxies and their central black holes also coevolve, although at a different rate from their counterparts in the field. A simple comparison between the AGN fraction in clusters and the field indicates that the cluster AGN fraction is approximately an order of magnitude lower over the entire observed redshift range.  One future direction for this research is to push further into the past toward the epoch of cluster galaxy assembly at $z\\sim2 - 3$. Numerous lines of evidence suggest that cluster galaxies formed earlier than their field counterparts, and the hypothesis of black hole and galaxy coevolution predicts that the the AGN fraction in dense environments should exceed the field value by this point. At lower redshifts, the identification of the properties of the ``transition'' group environment, in which the AGN fraction drops from the field to the cluster value, can potentially shed substantial light on the nature of the mechanism(s) responsible for triggering and fueling AGN. Finally, measurements of star formation and black hole accretion rates in the same cluster samples could reveal the extent of the correlation exhibited by individual galaxies and provide new insights into the physical processes that drive the observed coevolution.   \\begin{theacknowledgments} I am grateful for the contributions of T. Arnold, A. Berti, T.~E. Jeltema, D.~D. Kelson, J.~S. Mulchaey, G.~R. Sivakoff, and A.~I. Zabludoff to the research that I have presented in this conference article. I also appreciate the opportunity to speak at the conference and thank the organizers for this stimulating meeting. My research was supported in part by NASA through Chandra Award Number AR8-9014X and Spitzer Grant 1364997, and NSF via award AST-0705170, and the Department of Astronomy at OSU.  \\end{theacknowledgments}"
    },
    "0910/0910.2590.txt": {
        "abstract": "In this paper we present the results from the analysis of a sample of 28 gamma-ray burst (GRB) afterglow spectral energy distributions, spanning the X-ray through to near-infrared wavelengths. This is the largest sample of GRB afterglow spectral energy distributions thus far studied, providing a strong handle on the optical depth distribution of soft X-ray absorption and dust-extinction systems in GRB host galaxies. We detect an absorption system within the GRB host galaxy in 79\\% of the sample, and an extinction system in 71\\% of the sample, and find the Small Magellanic Cloud (SMC) extinction law to provide an acceptable fit to the host galaxy extinction profile for the majority of cases, consistent with previous findings. The range in the soft X-ray absorption to dust-extinction ratio, \\nhx/\\av, in GRB host galaxies spans almost two orders of magnitude, and the typical ratios are significantly larger than those of the Magellanic Clouds or Milky Way. Although dust destruction could be a cause, at least in part, for the large \\nhx/\\av\\ ratios, the good fit provided by the SMC extinction law for the majority of our sample suggests that there is an abundance of small dust grains in the GRB environment, which we would expect to have been destroyed if dust destruction were responsible for the large \\nhx/\\av\\ ratios. Instead, our analysis suggests that the distribution of \\nhx/\\av\\ in GRB host galaxies may be mostly intrinsic to these galaxies, and this is further substantiated by evidence for a strong negative correlation between \\nhx/\\av\\ and metallicity for a subsample of GRB hosts with known metallicity. Furthermore, we find the \\nhx/\\av\\ ratio and metallicity for this subsample of GRBs to be comparable to the relation found in other more metal-rich galaxies. ",
        "introduction": "To unravel the properties of gamma-ray burst (GRB) progenitors and the fundamental conditions required within a galaxy to form a GRB, an understanding of the GRB circumburst and host galaxy environment is essential. The faint optical magnitudes and large distances of GRB hosts limit the amount of information obtained from host galaxy observations, and thus studying the spectral properties of the afterglow provide the most direct and sensitive method of probing the surrounding environments of GRBs.  Both optical and X-ray spectroscopic observations have provided detailed information on the metal abundances and column densities in GRB local environments \\citep[e.g.][]{pcd+07,sff03, sf04}, and broadband analysis of the GRB afterglow spectral energy distribution (SED) allows the host galaxy dust extinction curve to be well modelled, and thus provides a measure of the host visual extinction. \\citet{gw01} combined these two techniques to compare the soft X-ray absorption with the visual dust extinction in the local environment of a sample of eight GRBs. In their analysis, they found that whereas the distribution of equivalent neutral hydrogen column density within GRB host galaxies was comparable to that observed in Galactic molecular clouds, the measured host galaxy visual extinction was 10--100 times smaller than expected for GRBs embedded in Galactic-like molecular clouds. \\citet{sfa+04} expanded on this work, and found the gas column density to dust extinction ratio to not only be larger than that of the Milky Way, but also to be around an order of magnitude larger than that of the gas rich Large and Small Magellanic Clouds (LMC and SMC, respectively). In both \\citet{gw01} and \\citet{sfa+04} the large column density to visual extinction ratios measured in GRB local environments was taken to be evidence of dust destruction by the GRB, causing the visual extinction to decrease. In addition to this, \\citet{sfa+04} also found that the optical to near-infrared (NIR) GRB afterglow data showed little evidence of the strong 2175~\\AA\\ absorption feature present in the Milky Way extinction law, and less prominently in the LMC. Instead, they found the mean SMC or the \\citet{dks+94} starburst galaxy dust extinction law to provide the best fit to their sample of GRB SEDs, both of which have no 2175~\\AA\\ absorption feature. These results are supported by the more recent work done by \\citet{kkz06} and \\cite{sww+07}. In each of these cases the mean SMC, LMC and Galactic extinction curves were assumed. However, a range in the total-to-selective extinction, $R_V=A_V/E(B-V)$, which relates the reddening to extinction, and 2175~\\AA\\ bump strength is observed along different lines-of-sight through the Milky Way \\citep{ccm89} and the Magellanic Clouds \\citep{gcm+03}, and observations of higher redshift supernovae and quasars ($z > 5$) also suggest differences between the dust extinction properties in higher redshift galaxies and the local universe \\citep[e.g.][]{mso+04,td01}. However, the degeneracy that exists between the best-fit GRB spectral index and the host galaxy's total-to-selective extinction means that to accurately determine the host galaxy's extinction law properties, good quality, broadband data are needed, preferentially stretching out into the negligibly extinguished far infrared (FIR) wavelength bands.  In the current era of \\swift\\ and rapid-response ground-based telescopes, prompt arcsecond GRB positions have provided a wealth of high quality, early-time X-ray, ultraviolet (UV), optical and NIR data. Accurate soft X-ray absorption measurements are now available for the large fraction of GRBs \\citep{crc+06,bk07,gnv+07,smp+07,ebp+09}, and well-sampled, high signal-to noise SEDs are providing strong constraints on the best-fit extinction law models \\citep[e.g.][]{pbb+08}. There have now been some examples of GRB host galaxies with the 2175~\\AA\\ absorption feature (e.g. GRB~070802, El{\\'{\\i}}asd{\\'o}ttir et al. 2009, Kr{\\\"u}hler et al. 2008; GRB~080607, Prochaska et al. 2009), as well as GRB host galaxies with \\Rv\\ values larger than the mean SMC, LMC and MW values \\citep[e.g.][]{pbb+08}, a possible indicator of grey dust, as suggested to be present in some GRB host galaxies \\citep[e.g.][]{sff03}. However, such analysis on the detailed properties of GRB extinction curves are typically still only possible for a handful of well-sampled, bright GRBs (e.g. GRB050525A, Heng et al. 2008; GRB~061126, Perley et al. 2009; GRB~070802, El{\\'{\\i}}asd{\\'o}ttir et al. 2009, Kr{\\\"u}hler et al. 2008; GRB~080607, Prochaska et al. 2009).  In \\citet{smp+07} we used X-ray and UV/optical simultaneous observations taken with the X-ray Telescope \\citep[XRT;][]{bhn+05} and Ultraviolet and Optical Telescope \\citep[UVOT;][]{rkm+05} onboard \\swift\\ \\citep{gcg+04} to analyse the SEDs for a sample of 7 GRBs. The dust extinction in the GRB host galaxies was modelled on the mean SMC, the LMC and the Milky Way extinction curves using the parameterisations given in \\citet{pei92}, which cover a range in 2175~\\AA\\ bump strengths and \\Rv\\ values. The SMC and LMC extinction curves were found to provide the best-fit model for the majority of the sample, in agreement with previous studies \\citep[e.g.][]{sfa+04,kkz06,sww+07}. However, we also found that, although the gas-to-dust ratio in \\swift\\ GRB host galaxies were typically larger than those of the Milky Way and Magellanic Clouds, the weighted mean was within 90\\% confidence of the Magellanic Clouds and Milky Way X-ray absorption to optical extinction ratios.  In this paper we aim to further our previous work, using a larger sample of 28 GRBs, and increasing the wavelength range of the afterglow SEDs to better constrain the absorption and extinction within the GRB host galaxy. In \\citet{smp+07} \\swift\\ data alone were used to produce the SEDs, and UVOT data with a rest-frame wavelength $\\lambda < 1215$~\\AA\\ were not included in the SED fits in order to avoid the absorption caused by the Lyman forest being confused for dust-extinction. In this paper we now model the absorption resulting from the Lyman forest such that all rest frame UV data redward of the Lyman edge is included in our spectral analysis. Furthermore, we also include additional ground-based NIR data if available, further increasing the spectral range of the SEDs and the degrees of freedom of the spectral fits. This provides better sampled SEDs and extends the redshift range within our sample, which was previously restricted to $z<1.7$ to ensure that the SED modelling was sufficiently well constrained within the optical wavelength range.  In $\\S$\\ref{sec:analysis} we present the new, extended GRB sample and describe the X-ray, UV/optical and NIR data reduction and analysis, and in $\\S$\\ref{sec:model} we describe the models used to fit the data. We present the results of our spectral modelling in $\\S$\\ref{sec:results} followed by an analysis of the possible selection effects and systematic biases that may be present in our work in $\\S$\\ref{sec:seleffs}. A discussion on the implications of our findings is presented in $\\S$\\ref{sec:disc}, and our conclusions are summarised in $\\S$\\ref{sec:cons}. Throughout the paper temporal and spectral indices, $\\alpha$ and $\\beta$, respectively, are denoted such that $F(\\nu,t)\\propto \\nu^{-\\beta}t^{-\\alpha}$, and all errors are $1\\sigma$ unless specified otherwise.   %====================================TABLE 1==================================== %=============================================================================== \\begin{table*} \\caption{Table listing the 28 GRBs in our sample with their redshift, Galactic column density and visual extinction in the line-of-sight to the GRB, the corresponding SED epoch, the UV, optical, and NIR band passes included in the GRB afterglow SED, and the rest frame coverage of the SED.\\label{tab:sample}} \\begin{tabular}{@{}lllllll} \\hline GRB & z & $N_{H,X}$(Gal) &  $A_V$(Gal) & Epoch & UV/optical/NIR bandpasses & Restframe Band \\\\ & & ($10^{21}$~\\invsqrcm) & & (s) & & Coverage (\\AA) \\\\ \\hline\\hline 050318 & 1.44$^a$ & 0.28 & 0.05 & T$^\\dag$+3600 & v,b,u & 1260--2400 \\\\ 050319 & 3.24$^b$ & 0.11 & 0.03 & T+20,000 & $I^{1,2},R^{1-3}$,v,b & 920--2090 \\\\ 050525A & 0.606$^c$ & 0.91 & 0.29 & T+20,000 & $K^{4,5},H^4,J^{4,6},I^{4,8},R^{4,9,10}$,v,b,u,w1,m2,w2 & 1000-14540 \\\\ 050730 & 3.968$^d$ & 0.30 & 0.16 & T+10,000 & $K^{11},J^{11},I^{11},R^{11}$,v,b & 780--4700 \\\\ 050802 & 1.71$^e$ & 0.18 & 0.06 & T+20,000 & $I^{12},R^{12}$,v,b,u,w1,m2,w2 & 590--3270 \\\\ 050820A & 2.6147$^f$ & 0.47 & 0.14& T+10,000 & $J^{13},z^{14},I^{14},R^{14},g^{14}$,v,b,u,w1 & 620--3710 \\\\ 050922C & 2.198$^g$ & 0.54  & 0.32 & T+20,000 & $R^{15-20}$,v,b,u,w1,m2 & 620--2360 \\\\ 051109A & 2.346$^h$ & 1.61 & 0.59 & T+5000 & $K^{21},H^{21},J^{21},I^{22},R^{23,24}$,v,b,u,w1 & 670--6980 \\\\ 060124 & 2.296$^i$ & 0.92 & 0.42 & T+100,000 & $I^{25},R^{25}$,v,b & 1180--2690 \\\\ 060206 & 4.048$^j$ & 0.09 & 0.04 & T+10,000 & $K^{26},H^{26},J^{26},R^{27-29}$,v,b & 770--4630\\\\ 060418 & 1.49$^k$ & 0.92 & 0.69 & T+5000 & $K^{30},H^{30},J^{30},z^{30,31},I^{32},R^{33}$,v,b,u,w1,m2 & 800--9380 \\\\ 060502A & 1.51$^l$ & 0.30 & 0.10 & T+5000 & $R^{34}$,v,b,u,w1 & 900--3010 \\\\ 060512 & 0.4428$^m$ & 0.14 & 0.05 & T+10,000 & $Ks^{35},J^{36},R^{37,38}$,v,b,u & 2130--16180 \\\\ 060526 & 3.21$^n$ & 0.55 & 0.21 & T+20,000 & $J^{39},I^{39,40},R^{40-47}$,v,b & 930--3180 \\\\ 060605 & 3.711$^o$ & 0.51 & 0.15 & T+10,000 & $R^{48-53}$,v,b & 830--1600 \\\\ 060607A & 3.082$^p$ & 0.27 & 0.09 & T+10,000 & $H^{30},i^{54},r^{54},g^{54}$,v,b,u & 750--4200 \\\\ 060714 & 2.71$^q$ & 0.61 & 0.24 & T+5000 & $R^{55}$,v,b & 1050--2040 \\\\ 060729 & 0.54$^r$ & 0.48 & 0.17 & T+70,000 & $R^{56}$,v,b,u,w1,m2,w2 & 1040--4900 \\\\ 060904B & 0.703$^s$ & 1.21 & 0.53 & T+5000 & $K^{57},J^{57},I^{57 },R^{58,59}$,v,b,u,w1,m2,w2 & 940--13710 \\\\ 060908 & 2.43$^t$ & 0.27 & 0.09 & T+5000 & $R^{60,61}$,v,b,u,w1 & 660--2200 \\\\ 060912 & 0.937$^u$ & 0.42 & 0.16& T+1500 & v,b,u,w1,m2 & 1030--3020 \\\\ 061007 & 1.262$^v$ & 0.21 & 0.06 & T+600 & $i^{62},R^{62}$,v,b,u,w1,m2,w2 & 710--3850 \\\\ 061121 & 1.314$^w$ & 0.51 & 0.14 & T+10,000 & $I^{62-64},R^{65}$,v,b,u,w1,m2,w2 & 690--3830 \\\\ 061126 & 1.159$^x$ & 1.00 & 0.56 & T+2000 & $K^{66},J^{66},I^{66,67},R^{66,68,69}$,v,b,u,w1,m2 & 920--10820 \\\\ 070110 & 2.352$^y$ & 0.18  & 0.04 & T+10,000 & $R^{70}$,v,b,u & 920--2250 \\\\ 070318 & 0.836$^z$ & 0.25 & 0.05 & T+1500 & v,b,u,w1,m2,w2 & 870--3190 \\\\ 070411 & 2.954$^\\ddag$ & 2.63 & 0.88 & T+500 & $R^{71,72}$,v,b & 990--1900 \\\\ 070529 & 2.4996$^\\S$ & 1.90 & 0.93 & T+600 & v,b,u,w1 & 640--1670 \\\\ \\hline \\end{tabular} \\begin{list}{}{} \\item[] $^a$ \\citet{bm05}; $^b$ \\citet{fhj+05}; $^c$ \\citet{fcb+05}; $^d$ \\citet{sve+05}; $^e$ \\citet{fsj+05}; $^f$ \\citet{lve+05}; $^g$ \\citet{jfp+05a}; $^h$ \\citet{qhr+05}; $^i$ \\citet{pft+06}; $^j$ \\citet{pwf+06}; $^k$ \\citet{dfp+06} $^l$ \\citet{cpf+06}; $^m$ \\citet{bfk+06}; $^n$ \\citet{bg06}; $^o$ \\citet{fkk+06}; $^p$ \\citet{lvs+06}; $^q$ \\citet{jvf+06d}; $^r$ \\citet{tlj+06}; $^s$ \\citet{fdm+06}; $^t$ \\citet{rjt+06}; $^u$ \\citet{jlc+06c}; $^v$ \\citet{jft+06b}; $^w$ \\citet{bpc06}; $^x$ \\citet{pbb+08}; $^y$ \\citet{jmf+07}; $^z$ \\citet{jfa+07}; $^\\ddag$ \\citet{jmt+07}; $^\\S$ \\citet{bfc07} \\item[]$^\\dag$ T is time at which the BAT triggered on the GRB \\item[] $^1$ \\citet{huk+07}; $^2$ \\citet{krm07}; $^3$ \\citet{wvw+05}; $^4$ \\citet{cb05}; $^5$ \\citet{rg05}; $^6$ \\citet{fhs+05}; $^7$ \\citet{ffc+05}; $^8$ \\citet{ytk05}; $^9$ \\citet{hhg+05}; $^{10}$ \\citet{mbs05a}, $^{11}$ \\citet{pcm+06}; $^{12}$ \\citet{pes+05}; $^{13}$ \\citet{meb+05}; $^{14}$ \\citet{ckh+06}; $^{15}$ \\citet{dp05}; $^{16}$ \\citet{jpt+05b}; $^{17}$ \\citet{ap05}; $^{18}$ \\citet{hkh+05}; $^{19}$ \\citet{pmm+05}; $^{20}$ \\citet{dpf+05}; $^{21}$ \\citet{bbs+05}; $^{22}$ \\citet{tor05}; $^{23}$ \\citet{mwp+05}; $^{24}$ \\citet{juc+05}; $^{25}$ \\citet{mbs+07}; $^{26}$ \\citet{apb06}; $^{27}$ \\citet{cvw+07}; $^{28}$ \\citet{sdp+07}; $^{29}$ \\citet{wvw+06}; $^{30}$ \\citet{mvm+07}; $^{31}$ \\citet{nif+06}; $^{32}$ \\citet{cob06a}; $^{33}$ \\citet{kop06}; $^{34}$ \\citet{cof06}, $^{35}$ \\citet{hlm+06}; $^{36}$ \\citet{sdp+06}; $^{37}$ \\citet{cen06a}; $^{38}$ \\citet{mil06}; $^{39}$ \\citet{cob06b}; $^{40}$ \\citet{tgb+06}; $^{41}$ \\citet{bgv+06}; $^{42}$ \\citet{cig+06}, $^{43}$ \\citet{dhm+07}; $^{44}$ \\citet{gtb+06}; $^{45}$ \\citet{kbs+06a}; $^{46}$ \\citet{md06}; $^{47}$ \\citet{rpi+06} $^{48}$ \\citet{kg06}; $^{49}$ \\citet{ksa+06c}; $^{50}$ \\citet{ksa+06b}; $^{51}$ \\citet{mfm+06}; $^{52}$ \\citet{sap06}; $^{53}$ \\citet{zqw+06}; $^{54}$ \\citet{nrc+09}; $^{55}$ \\citet{api06}; $^{56}$ \\citet{qr06}; $^{57}$ \\citet{cb06}; $^{58}$ \\citet{pks+06}; $^{59}$ \\citet{skv06}; $^{60}$ \\citet{act+06}; $^{61}$ \\citet{wtr06}; $^{62}$ \\citet{mmg+07}; $^{62}$ \\citet{cen06b}; $^{63}$ \\citet{cob06c}; $^{64}$ \\citet{tor06a}; $^{65}$ \\citet{uau06}; $^{66}$ \\citet{pbb+08}; $^{67}$ \\citet{tor06b}; $^{68}$ \\citet{smg+06}; $^{69}$ \\citet{wm06}; $^{70}$ \\citet{mjv07}; $^{71}$ \\citet{msd07}; $^{72}$\\citet{klk+07} \\end{list} \\end{table*} %=============================================================================== ",
        "conclusions": "\\label{sec:cons} In this paper we have presented the results from the spectral analysis of 28 GRB SEDs. We measured the equivalent neutral hydrogen column density and visual extinction at the host galaxy, and found 79\\% of the GRBs in our sample to have a detectable soft X-ray absorption system in the host galaxy, and 71\\% to have a detectable visual dust-extinction system. Using the measured \\nhx/\\av\\ ratios as an indicator of the host galaxy gas-to-dust ratio, we find that GRB host galaxies have gas-to-dust ratios that are typically larger than those measured in the Milky Way and Magellanic Clouds by up to two orders of magnitude. We have investigated several possibilities that could account for the relatively large gas-to-dust ratios in GRB host galaxies.  There is no evidence to suggest that the large host galaxy \\nhx/\\av\\ ratios measured in our GRB sample is the result of any systematic error in the way that we measure \\av. One possibility is that dust destruction by the GRB has reduced the visual extinction, \\av, relative to the equivalent neutral hydrogen column density, \\nhx. However, there are currently no observations that clearly show the early time colour evolution expected from dust destruction. Although such observations are limited by the quality and promptness of the data, we also found that the majority of our sample had host dust properties best-fit by the UV steep, SMC extinction law, indicating an abundance of small dust grains in the GRB surrounding environment. In the event of a significant phase of dust destruction, a grey extinction law should be observed, where the differential change in extinction from UV to NIR energy bands is small. The dust probed by our \\av\\ measurements must, therefore, lie in regions of the GRB host galaxy that have not been subjected to significant amounts of dust destruction.  For a subset of eight GRBs we were also able to study how the neutral hydrogen column density, \\nh, compared with \\av, and we found \\nh/\\av\\ to extend to both larger and smaller values than those of the Magellanic Clouds and the Milky Way by up to an order of magnitude. The distribution in \\nh/\\av\\ can be accounted for by the competing effects that alter the values of \\nh\\ and \\av. Firstly, differences in the host galaxy metallicities and in the amount of dust destroyed by the GRB will affect the value of \\av. On the other hand, the value of \\nh\\ will be dependent on the amount of photo-ionised hydrogen along the line-of-sight to the GRB. The mean logarithmic metallicity of the GRB sample with both \\nh\\ and \\av\\ measurements is almost 1.0~dex smaller than that of the SMC (0.04~$Z_{\\odot}$), and we would therefore expect the GRB host galaxy \\nh/\\av\\ ratio to be significantly smaller than the SMC \\nh/\\av\\ ratio. The roughly even number of GRBs with smaller and larger \\nh/\\av\\ ratios than the Magellanic Clouds and Milky Way therefore implies that the level of photo-ionised hydrogen along the line-of-sight to the GRB is greater than the fraction of dust destroyed by the GRB. This would suggest that measurements of \\nhx\\ and \\av\\ probe regions of dust and gas much closer to the GRB than \\nh.  It has been suggested that differences in the gas-to-dust ratios in galaxies of different types are correlated with the metallicity of the galaxy \\citep[e.g.][]{ddb+07}, whereby smaller metallicity systems have larger gas-to-dust ratios. From a subsample of four GRBs with measured metallicity and a soft X-ray absorption and visual extinction system detected with $90$\\% confidence, together with the Small and Large Magellanic Clouds and Milky Way, we found a strong negative correlation between the \\nhx/\\av\\ ratio and the metallicity, [M/H]. The spearman rank coefficient was -0.89 with $90$\\% confidence. The large \\nhx/\\av\\ ratios measured in GRB host galaxies could, therefore, be an indication of their very low, although broad, range of metallicities. A greater sample of GRB hosts with measured metallicities are needed to verify such a correlation, which if confirmed would suggest that low-metallicity environments are less efficient at forming dust from their metals than high-metallicity galaxies."
    },
    "0910/0910.3655_arXiv.txt": {
        "abstract": "Observations of transient phenomena in the Universe reveal a spectrum of mass-ejection properties associated with massive stars, covering from Type II/Ib/Ic core-collapse supernovae (SNe) to giant eruptions of Luminous Blue Variables (LBV) and optical transients. In this work, we hypothesize that a large fraction of these phenomena may have an explosive origin, the distinguishing ingredient being the ratio of the prompt energy release $E_{\\rm dep}$ to the envelope binding energy $E_{\\rm binding}$. Using one-dimensional one-group radiation hydrodynamics and a set of 10-25\\,\\msun\\, massive-star models, we explore the dynamical response of a stellar envelope subject to a strong, sudden, and deeply-rooted energy release. Following energy deposition, a shock {\\it systematically} forms, crosses the progenitor envelope on a day time-scale, and breaks-out with a signal of hour-to-days duration and a 10$^5$-10$^{11}$\\,\\lsun\\, luminosity. We identify three different regimes, corresponding to a transition from dynamic to quasi-static diffusion transport. For $E_{\\rm dep} > E_{\\rm binding}$, full envelope ejection results with a SN-like bolometric luminosity and kinetic energy, modulations being commensurate to the energy deposited and echoing the diversity of Type II-Plateau SNe. For $E_{\\rm dep} \\sim E_{\\rm binding}$, partial envelope ejection results with a small expansion speed, and a more modest but year-long luminosity plateau, reminiscent of LBV eruptions or so-called SN impostors. For $E_{\\rm dep} < E_{\\rm binding}$, we obtain a ``puffed-up'' star, secularly relaxing back to thermal equilibrium. In parallel with gravitational collapse and Type II SNe, we argue that thermonuclear combustion, for example of as little as a few 0.01\\,\\msun\\, of C/O, could power a wide range of explosions/eruptions. Besides massive stars close to the Eddington limit and/or critical rotation, 8-12\\,\\msun\\, red-supergiants, which are amongst the least bound of all stars, represent attractive candidates for transient phenomena. ",
        "introduction": "\\label{sect_intro}  Stellar explosions, broadly refered to as supernovae (SNe), are understood to stem from a sudden release of energy either associated with the collapse of the degenerate core of a massive star or from the thermonuclear combustion of fresh fuel deep inside the stellar envelope. Whether one or the other mechanism occurs seems to depend on the main-sequence mass of the progenitor star, with core collapse occuring systematically if its value is above $\\sim$8\\,\\msun\\, \\citep[hereafter WHW02]{WHW02}. Interestingly, whatever the mechanism, the typical kinetic energy of SN ejecta is on the order of 10$^{51}$\\,erg, as inferred for example for the well-studied SN 1987A (Type II peculiar; \\citealt{blinnikov_etal_2000}), for SN 1999em (Type II-Plateau, heareafter II-P; \\citealt{utrobin_07}), or for the very uniform set of events that Type Ia SNe constitutes \\citep{WKB07_Ia_lc}. The cause of this apparent degeneracy in explosion energy is, paradoxically, perhaps not so much tied to the mechanism itself, but instead to the rather uniform total envelope binding energy of the progenitor stars, on the order of 10$^{51}$\\,erg; anything falling short of that leads to a fizzle and no SN display.  \\begin{figure*} \\epsfig{file=f1a.ps,width=8.5cm} \\epsfig{file=f1b.ps,width=8.5cm} \\caption{{\\it Left:} Comparison of absolute-$V$-band-magnitude light curves for an illustrative and non-exhaustive sample of SNe, SN impostors and/or erupting LBVs (violet: SN1999em, \\citealt{leonard_etal_02a,DH06_SN1999em}; blue: SN1999gi, \\citealt{leonard_etal_02b}; turquoise: SN 2005cs, \\citealt{pastorello_etal_09}; green: SN1999br, \\citealt{pastorello_etal_09}: light green: $\\eta$ Car, \\citealt{frew_04}; yellow: SN 1997bs, \\citealt{vandyk_etal_00}; red: NGC2363-V1, \\citealt{drissen_etal_01,petit_etal_06}). For each, we adopt the distance and reddening given in the associated references. The time origin is that of maximum recorded brightness. Note that all these objects have comparable effective temperatures on the order of 10,000\\,K, hence comparable bolometric correction, making the comparison of their absolute $V$-band magnitude meaningful. We also show on the left side and in black the absolute visual magnitude of the galactic red-supergiant stars studied by \\citet[black]{levesque_etal_05}, as well as the {\\it bolometric} magnitude of the O star models computed by \\citet[gray; we use the bolometric magnitude here since O stars are hot and have large bolometric corrections]{martins_etal_05}. {\\it Right:} Same as left, but now zooming in on the time of maximum brightness (the color coding is the same as in the left panel). Notice the stark contrast between SN light curves, associated with shorter/brighter events, and erupting massive stars, associated with longer/fainter events. Importantly, notice the overlap between the intrinsic brightness of $\\eta$ Car and that of the low-luminosity Type II-P SN 1999br. In this work, we propose that this diversity of radiative displays may be accomodated by a {\\it common, explosive, origin}. \\label{fig_obs_lc} } \\end{figure*}  The last decade of observations of such transient phenomena has shown, however, that the radiative signatures associated with SNe (as classified in circulars) are very diverse, from very faint to very luminous, from fast-evolving to slow-evolving or fast-expanding to slow-expanding. This diversity has been observed little in Type Ia SNe, with a few peculiar events such as SN 2002ic (presence of narrow hydrogen lines in an otherwise standard Type Ia spectrum; \\citealt{hamuy_iau_02ic}) or SNLS-03D3bb (possible Type Ia SN from a super-Chandrasekhar white dwarf star; \\citealt{howell_etal_06}). In contrast, there has been a rich diversity in explosions associated (perhaps erroneously at times) with massive stars and the mechanism of core collapse. We have observed 1) Type Ic SNe, associated or not with a long-soft $\\gamma$-ray burst, and with a standard or a very large kinetic energy (SN 1998bw, \\citealt{WES_99}; SN 2002ap, \\citealt{mazzali_etal_02}); 2) a large population of Type II-Plateau (II-P) SNe in what seems to be the generic explosion of a moderate-mass red-supergiant (RSG) star (e.g. SN 2005cs, \\citealt{maund_etal_05,UC_08}); 3) a growing number of low-luminosity SNe that share properties with standard Type II-P SNe except for being significantly and globally less energetic (e.g. SN1997D, \\citealt{chugai_utrobin_00}; SN 1999br, \\citealt{pastorello_etal_04}; OT2006-1 in M85, whose status is ambiguous, see \\citealt{kulkarni_etal_2007_m85,pastorello_etal_07_m85}). We show a sample of $V$-band absolute-magnitude light curves of such core-collapse SNe in Fig.~\\ref{fig_obs_lc}, with representative peak values  of $-$14 to $-$17\\,mag and a 100-day plateau duration (best seen in the right panel of that figure), hence about 6-10\\,mag brighter than their proposed RSG progenitors (shown as black crosses). From this expanded SN sample, the range of corresponding explosion energies has considerably widened, extending above and below the standard 10$^{51}$\\,erg value. Within the core-collapse SN context, this modulation is thought to stem from modulations in the energy revival of the stalled shock above the nascent proto-neutron star, in turn modulated by the stellar-core structure \\citep{burrows_etal_07a}.  The stretching to low explosion energies of potential core-collapse SN events is intriguing. For the most energetic explosions belonging to points 1 and 2 above, the classification as a SN is unambiguous. However, some transient events show an ejecta/outflow kinetic energy and a peak magnitude that are SN-like, although the events did not stem from core collapse (a star is observed at that location on post-explosion/eruption images); the community calls these SN impostors (e.g. SN1997bs; Fig.~\\ref{fig_obs_lc}; \\citealt{vandyk_etal_00}). Conversely, this raises the issue whether {\\it low-energy} Type II-P SNe are associated with core collapse - they might but they need not. We illustrate this overlap in radiative properties for a sample of such objects in Fig.~\\ref{fig_obs_lc}. One such case is the Luminous Blue Variable (LBV) $\\eta$ Car, whose properties during its 1843 eruption rival those of the low-luminosity Type II-P SN1999br. $\\eta$ Car survived this gigantic eruption, which shed about 10\\,\\msun\\, of material in what now constitutes the homunculus nebula \\citep{smith_etal_2003}. In contrast to core-collapse SNe, such eruptive phenomena in massive stars have been associated with the proximity of the star to the Eddington luminosity $L_{\\rm Edd} = 4 \\pi c G M / \\kappa$ ($\\kappa$ is the mass absorption coefficient). Due to the steep dependence of luminosity to mass (e.g. with an exponent of 3.5 for main-sequence objects), this limit is easily reached by very massive stars such as $\\eta$ Car, or more generally massive blue-supergiant stars. In this context, massive stars are thought to undergo considerable mass loss when their luminosity overcomes the Eddington limit,\\footnote{Note, however, that energy will have to be supplied to the stellar envelope to push it over this limit, and in large amounts to explain such a nebula as the homunculus.} giving rise to a porosity-modulated continuum-driven outflow \\citep{shaviv_00,owocki_etal_04}. Here, this super-Eddington wind constitutes a quasi steady-state outflow, and has therefore been thought to be of a fundamentally different nature from core-collapse SN ejecta. And indeed, one refers to a wind for the former and to an ejecta for the later. This dichotomy has been exacerbated by the stark contrast in typical light curves of eruptive stars (long lived with large brightness) and core-collapse SN explosions (short lived with huge brightness). In Fig.~\\ref{fig_obs_lc}, we show two known eruptive massive stars that highlight this contrast.  However, recent observations may be challenging such a strict segregation. First, the recent identification of very fast outflowing material ahead of $\\eta$ Car's homunculus now suggests that such material was accelerated by a shock, rather than driven in a quasi-steady wind, and thus connects the giant outburst to an explosive origin \\citep{smith_08_blast}. Second,  the existence of interacting SNe tells us that a massive eruption can occur merely a few years before explosion. For some, e.g. SN 2006gy \\citep{smith_etal_07a,smith_etal_07b, woosley_etal_07} or SN 1994W \\citep{dessart_etal_09}, the amount is thought to be large enough to decelerate the energetic (and necessarily faster-expanding) subsequent ejection. This very strict timing of merely a few years, {\\it which is orders of magnitude smaller than evolutionary or transport time-scales}, suggests a connection between the mechanisms at the origin of the two ejections. For SN2006gy, Woosley et al. propose recurrent pair-instability pulsations, a mechanism germane to super-massive stars and therefore extremely rare. For lower mass massive stars, this short delay of a few years seems to exclude a very-long, secular, evolution for the production of the first ejection since this would have no natural timing to the comparatively instantaneous event of core collapse. Motivated by these recent observations, we explore in this paper whether this diversity of events could be reproduced by a unique and deeply-rooted mechanism, associated with the sudden energy release above the stellar core and the subsequent shock heating of the progenitor envelope. This means would communicate a large energy to the stellar envelope on a shock crossing time-scale of days rather than on a very long-diffusion time-scale of thousands of years or more. Although different in their origin, this energy release could be a weak analogue of what results in pair-instability pulsations, i.e. a nuclear flash, as identified in the 8-12\\,\\msun\\, range by \\citet{weaver_woosley_79}.  In this paper, following this shock-heating hypothesis, we use 1D radiation-hydrodynamics simulations to explore the production of explosions/eruptions in stars more massive that $\\sim$10\\,\\msun\\, on the main sequence. Rather than focusing on specific models, like those potentially associated with failed supernovae \\citep{fryer_etal_09}, we parameterize the problem through a simple energy deposition, taking place with a given magnitude, over a given time, and at a given depth in a set of pre-SN progenitor star models. We do not aim at reproducing any specific observation but, through a systematic approach, try to identify important trends, in a spirit similar to that of \\citet{falk_arnett_77}. However, we depart from these authors by studying ``non-standard''  explosions. In practice, we consider cases where the energy deposited can be both smaller or larger than the binding energy of the overlying envelope, but must imperatively be released on a very short time-scale to trigger the formation of a shock. Doing so, we identify three regimes, with ``standard'' SN explosions (short-lived transients) at the high energy end, objects that we will group in the category SN ``impostors'' (long-lived transients) at intermediate energy, and variable stars at the very low energy end. Let us stress here that we do not make the claim that all massive-star eruptions, or all transients in general, stem from a strong, sudden, and deeply-rooted energy release in their envelope. Here, we make this our working hypothesis and investigate how much such an explosive scenario can explain observations. We do find that this scenario has great potential and should be considered as a possibility when examining the origin of massive-star eruptions and associated transient phenomena.  \\begin{figure} \\epsfig{file=f2.ps,width=8.5cm} \\caption{Density distribution as a function of Lagrangian mass at the onset of collapse for the models of WHW02 evolved at solar metallicity. Notice the flattening density distribution above the core for increasing mass. In low-mass massive stars, the star is structured as a dense inner region (the core), and a tenuous extended H-rich envelope. \\label{fig_rho_mr} } \\end{figure}  The paper is structured as follows. In \\S\\ref{sect_input}, we briefly present the stellar evolutionary models of WHW02 that we use as input for our 1D 1-group radiation-hydrodynamics simulations. We discuss the properties of the progenitor massive stars of WHW02, such as density structure and binding energy, that are relevant for the present study. We then describe in \\S\\ref{sect_model}, the numerical technique and setup for our energy deposition study. In \\S\\ref{sect_s11}, we present the results of a sequence of simulations based primarily on the 11\\,\\msun\\, model of WHW02, discussing the properties of the shocked progenitor envelope for different values of the strength (\\S\\ref{var_edep}), the depth (\\S\\ref{var_mcut}) and the duration (\\S\\ref{var_dt}) of the energy deposition. We also discuss in \\S\\ref{var_mprog} the results obtained for more massive pre-SN progenitors, ranging from 15 to 25\\,\\msun\\, on the main sequence. In \\S\\ref{sect_rad}, we present synthetic spectra computed separately with CMFGEN \\citep{HM98_lb,DH05_qs_SN,dessart_etal_09} for a few models at a representative time after shock breakout. In \\S\\ref{sect_discussion}, we discuss the implications of our results for understanding transient phenomena, and in \\S\\ref{sect_conclusion} we summarize our conclusions. ",
        "conclusions": "\\label{sect_conclusion}  In this paper, we have presented one-dimensional one-group (gray) two-temperature radiation-hydrodynamics simulations of pre-SN massive-star envelopes subject to a sudden release of energy above their degenerate core. Although at the high energy end, the likely phenomenon is core collapse, we more generally have in mind the thermonuclear incineration of intermediate-mass elements present in shells above the core. The motivation for this study is: 1) The existence of interacting SNe (which must eject a shell at least once before core-collapse, but more likely eject multiple shells, by some means not elucidated so far); 2) the identification of fast-moving material exterior to $\\eta$ Car's homunculus (which cannot stem from radiation-driving in a wind;  \\citealt{smith_08_blast}); 3) the broad range of energies inferred for Type II-P SNe, overlapping at the low-energy end with high-energy transients and massive-star eruptions. The bulk of our results stem from work on the 11\\,\\msun\\, model of WHW02, although tests with higher-mass progenitors yield the same qualitative conclusions and regimes, merely shifted quantitatively. This work is an exploration of the outcome of a strong, sudden, and deeply-rooted energy deposition taking place at the base of a massive-star envelope. We are not proposing this scenario holds for all massive-star mass ejections, but we investigate what results when this circumstance obtains. There is no doubt it does, but how frequently and how robustly is left for future study.  Although this result is not new, we find that the fundamental quantity controling the outcome of a strong, sudden, and deeply-rooted energy deposition is the binding energy of the stellar envelope, which increases considerably with progenitor mass in the range 10--40\\,\\msun. What is new, however, is our study of the long-term evolution of the gas and radiative properties that result from configurations where the energy deposited is greater, on the order of, or less than the envelope binding energy. We identify three regimes, with a continuous progression from 1) SN explosions at high energy ($E_{\\rm dep} > E_{\\rm binding}$), with a complete envelope ejection associated with a 100-day long high-plateau luminosity; 2) SN impostors at the intermediate energy range ($E_{\\rm dep} \\sim E_{\\rm binding}$), with a partial envelope ejection, and a more modest but longer-lived plateau luminosity; and 3) bloated/variable stars at the low-energy end ($E_{\\rm dep} < E_{\\rm binding}$), with little or no mass ejection but with a residual puffed-up envelope of weakly-modified properties. What conditions the results is not the magnitude of the energy deposition itself but how it compares with the binding energy of the envelope exterior to the site of deposition. Hence, to achieve the same result requires more energy in a more massive progenitor star. These properties are summarized in Fig.~\\ref{fig_summary_ekin}.  \\begin{figure} \\epsfig{file=f14.ps,width=8.5cm} \\caption{Correlation between various energies (normalized to the corresponding adopted energy deposition) and the energy deposition (normalized to the corresponding envelope binding energy). In black, we show the normalized asymptotic energy of the ejecta/outflow, and in blue (red) the normalized envelope/ejecta internal (kinetic) energy at the time of shock breakout. Only the models presented in Table~1 are shown here. The black curve illustrates the three regimes we have presented here, with SN explosions, SN impostors, and variable/perturbed stars as the energy deposition varies from being much larger, comparable, and smaller than the envelope binding energy. Notice how the internal energy always dominates over the kinetic energy at the time of shock breakout in the simulations performed here. \\label{fig_summary_ekin} } \\end{figure}    In all simulations presented, a shock forms systematically with a strength that depends on the energy deposited. It crosses the envelope in  $\\sim$1 to $\\sim$50 days, hence communicating its energy to the entire progenitor envelope {\\it quasi-instantaneously}, i.e. as opposed to, e.g., a diffusion time-scale for energy transport of 10$^4$ years or more at the corresponding depth. This shock eventually emerges at the progenitor surface with a breakout signal that varies from a duration of an hour up to a few days (modulated here by the shock speed rather than by the atmospheric-scale height), with a flux peaking at longer wavelengths for weaker shock strengths. This breakout signal is the (and may be the only) {\\it unambiguous} evidence that the subsequent ``ejection\" was triggered by shock-heating, and thus has an explosive origin.  At shock breakout, the luminosity reaches a peak, then fades, before stabilizing/rising again forming an extended plateau. This plateau phase corresponds to the recombination epoch of the ejected mass, the internal energy decreasing primarily through expansion and little through radiation. It is the large stored internal energy (in the form of radiation) that keeps the ejecta optically thick and luminous (radioactive decay is neglected in our work). We find a continuum of light curves, faster-expanding ejecta being more luminous both at breakout (stronger shock) and during the plateau phase (Fig.~\\ref{fig_summary_lpeak}). The models presented in \\S\\ref{var_edep} corroborate the correlation between plateau luminosity and  mid-plateau photospheric velocity identified by \\citet{hamuy_pinto_02}, and refined by \\citet{nugent_etal_06,poznanski_etal_09}. We also find that the plateau duration is anti-correlated with energy deposition (Fig.~\\ref{fig_summary_dt}). At larger energy, faster expansion leads to faster cooling and recombination so that the ejecta photosphere recedes faster in mass/radius after reaching its peak earlier. For small energy variations in this regime, interplay between kinetic and internal energy (which are comparable at breakout) yield a plateau duration that is $\\sim$100\\,d, which is on the order of Type II-P plateau lengths. For lower energy deposition, we switch slowly from a regime of dynamic diffusion to that of quasi-static diffusion. The more slowly-expanding ejecta, characterized by a slowly-decreasing optical depth with time, gives the bolometric luminosity a modest peak value and a slow evolution, with plateau durations of up to 1-2 years.  The plateau phase is thus more extended and fainter for lower energy deposition, echoing the light-curve trend going from SNe to SN impostors (Fig.~\\ref{fig_obs_lc}). Note that, in our simulations, these time-scales are always 6-7 orders of magnitude larger than the time-scale for the energy deposition, which was chosen to be ten seconds in most cases. In other words, the time-scale over which the light curve evolves (basically that of radiative diffusion in an expanding medium) has no connection to the time-scale over which the energy was deposited in the first place.  \\begin{figure} \\epsfig{file=f15.ps,width=8.5cm} \\caption{Correlation between the peak luminosity during the plateau phase (black dots) and the mass-weighted average ejecta velocity. We also show the correlation for the peak luminosity at shock breakout (red dots; scaled down by a factor of 1000 for convenience). Note that we include models from Tables~1 and 2. We overplot the line $L \\propto v^{1.6}$, which follows closely the distribution of points for $L_{\\rm Peak, Plateau}$ versus $\\langle v \\rangle_M$. Our radiation-hydrodynamics simulations support the correlation identified by \\citet{hamuy_pinto_02} and subsequently improved upon by \\citet{nugent_etal_06,poznanski_etal_09}. Our slope is in close agreement with that proposed in this last reference. Impressively, the relation holds over the entire domain explored. Note that there is no consideration of radioactive decay from unstable isotopes or departures from spherical symmetry, and only data points associated with the 11\\,\\msun-progenitor star are used. Relaxing these choices would likely introduce some scatter. \\label{fig_summary_lpeak} } \\end{figure}  \\begin{figure} \\epsfig{file=f16.ps,width=8.5cm} \\caption{Correlation between the plateau duration (black) and the time-like quantity $R_{\\rm phot,P}/V_{\\rm phot,P}$ (ratio of the radius and the velocity at the photosphere at the time of peak-plateau brightness) versus the time-like quantity equal to the asymptotic ejecta kinetic energy divided by the peak-plateau luminosity brightness. \\label{fig_summary_dt} } \\end{figure}  From our exploration and with our set of models, we find that explosions of varying strength can yield the broadest range of outcomes in {\\it low-mass} massive stars because they are characterized by a very low envelope binding energy (Fig.~\\ref{fig_eb_mr}). We indeed obtain light curves evolving from week to year time-scales (Fig.~\\ref{fig_s11_lum_all}) and ejecta expansion rates ranging from a few tens to a few thousand \\kms (Fig.~\\ref{fig_s11_phot}). An explosion/eruption producing a transient requires merely 10$^{49}$\\,erg in such objects, a value that is so low that gravitational collapse of the stellar core may not be required. And indeed, in this mass range, stellar-evolutionary calculations reveal the existence of nuclear flashes in the last nuclear-burning stages \\citep{weaver_woosley_79}, which could represent such an energy source. We therefore propose that low-mass massive stars are prime candidates for transient phenomena in the Universe, as well as prime canditates for interacting SNe such as 1994W \\citep{dessart_etal_09}.  In our simulations, sudden energy deposition above the core leads to shock-heating of the {\\it entire} envelope. Whenever the energy deposited is greater than its binding energy, the {\\it entire} envelope is ejected, with an asymptotic kinetic energy that is commensurate with the energy deposited. If a subsequent energy deposition occurs (e.g. a nuclear flash followed by gravitational collapse, as needed in a fraction at least of interacting SNe), the second ejection would have low mass and little or no hydrogen. Depending on the time between the two ejections, one could get an interacting SN for a short delay (i.e. a few years) or two transients at the same location for a long delay (the first one being dim if the explosion energy is small and the second one being potentially very dim due to the low ejected mass). These scenarios are rather speculative, but they are warranted since at present most, if not all, transients have an unknown origin and are poorly characterized.  In our simulations, and within the context of this work, we obtain small mass ejections only when depositing the energy close to the progenitor surface, an eventuality that seems difficult to justify physically in such massive stars. Such low-mass ejections would seem to be better suited, for example, to the surface layers of an extended white dwarf. Observationally, low-mass ejections are likely to be associated with fast transients. For transients that are both fast and faint, a low-mass ejection in a highly- or moderately-bound object seems required. At the least, our simulations for (loosely-bound) RSG stars perturbed by a small deeply-rooted energy release produce large mass ejections (the overlying hydrogen envelope) and long-faint transients.  Synthetic spectra computed for a sequence of models with varying energy deposition reveal a continuous evolution from Type II-P SN-like spectra at high energy ($L\\sim$10$^8$\\,\\lsun), to low-luminosity Type II-P SN spectra at intermediate energy ($L\\sim$10$^7$\\,\\lsun), to SN-impostor-like spectra at low energy  ($L\\sim$10$^6$\\,\\lsun), with, in the same order, narrower line profiles and redder/cooler spectral-energy distributions (Fig.~\\ref{fig_v1d_sed}).  The results from this work should not be compromised by the approximations made in our radiation-hydrodynamics simulations. First, with one dimensionality we prevent convective transport and any structure formation through, e.g., Rayleigh-Taylor instabilities. This may alter the properties of models characterized by low expansion speeds (longer evolution time-scale); we thus plan to study this eventuality with 2D and 3D simulations in the future. Second, we deposit the energy at a large and fixed rate, independent of any feedback effects. In the case of nuclear flashes, such feedback effects could shut-off the burning prematurely. We are aware of this artifact and will attempt in future work to develop a more physically-consistent approach by investigating the conditions that may lead to shock formation, rather than assuming a setup that systematically leads to it. However, provided a shock forms, we think our results apply. Third, progenitor models may differ on a mass-by-mass comparison with other groups but the general trend of increasing binding energy with main sequence mass should hold. Fourth, one-group transport should be accurate enough since it has been shown to capture the essentials of such radiation hydrodynamics simulations \\citep{utrobin_chugai_09} - the key physics presented here takes place at large optical depth, under LTE conditions.  Our finding that very modest energy perturbations can dramatically affect the structure of a massive star motivates detailed multi-dimensional hydrodynamical investigations of massive-star interiors, in particular of the last burning stages preceding core-collapse (see, e.g., \\citealt{bazan_arnett_94,bazan_arnett_98,asida_arnett_00, arnett_etal_05,meakin_arnett_06}). As shown in these hydrodynamical simulations, and more generally, we surmise that the quasi-steady state approach of standard stellar-evolutionary codes (which keep an eye on the longer-term evolution of stars) may be missing important ingredients for our understanding of massive-star evolution. This has relevance for understanding mass loss and in particular massive-star eruptions, stellar variability/stability, and interacting SNe. Such pre-SN mass ejections would also modify, and perhaps sizably, the envelope mass and structure, thereby affecting the conditions surrounding core collapse and explosion.  Ongoing and forthcoming surveys/missions like Pan-STARRS, Sky Mapper, the Palomar Transient Factory, GALEX, or the Large Synoptic Survey Telescope will better reveal the diversity of explosions, eruptions, and more generally transient phenomena in the Universe. We surmise that surveys up to now have missed a large number of low-energy long-lived transients, such as low-luminosity Type II-P SNe (objects even less energetic than SN1999br) and SN impostors. It is somewhat surprising that we have not yet detected Type II SNe with Plateau durations well in excess of 100 days. Moreover, for the shock-heating solutions presented here, a breakout signal systematically takes place. At SN-like energies, the signal may be too short to be resolved \\citep{gezari_etal_08}, but for lower-energy transients, the reduced shock speed and strength would lengthen the breakout duration up to about a day, and move the peak of the spectral-energy distribution from $\\sim$100\\AA\\ to the 1000-3000\\AA\\ range. Hence, the breakout signal should be more easily detectable in such transients, allowing to distinguish between an explosive event and, e.g., a super-Eddington wind."
    },
    "0910/0910.0261_arXiv.txt": {
        "abstract": "We have been analyzing a large sample of solar-like stars with and without planets in order to homogeneously measure their photospheric parameters and Carbon abundances. Our sample contains around 200 stars in the solar neighborhood observed with the ELODIE spectrograph, for which the observational data are publicly available. We performed spectral synthesis of prominent bands of C$_{2}$ and C~I lines, aiming to accurately obtain the C abundances. We intend to contribute homogeneous results to studies that compare elemental abundances in stars with and without known planets. New arguments will be brought forward to the discussion of possible chemical anomalies that have been suggested in the literature, leading us to a better understanding of the planetary formation process. In this work we focus on the C abundances in both stellar groups of our sample. ",
        "introduction": "The Sun is widely thought to be formed from material representative of local physical conditions in the Galaxy at the time of its formation and to represent a standard chemical composition. This {\\it homogeneity hypothesis} has been often put in question because of many improvements in the observations techniques and data analysis. With the discovery of new planetary systems, the already known {\\it heterogeneity} sources (stellar nucleosynthesis, stellar formation process) have gained a new perspective and brought new questions. It is now a fact that stars with giant planets are, on average, richer in metals than those for which no planet was detected (\\cite[Gonzalez 2006]{Gonzalez2006} and references therein). Some authors suggested that this kind of anomaly may not only involve the content of heavy elements but also some light elements like Li, C, N, and O (\\cite[Ecuvillon et al. 2006]{Ecuvillonetal2006} and references therein). Other authors found no difference in the abundance of light elements when comparing stars with and without planets (see \\cite[Ecuvillon et al. 2004]{Ecuvillonetal2004} and references therein). The studies above are not conclusive yet. We need new tests, using more accurate and homogeneous data, with a larger number of stars. Abundances of these elements in solar-like stars will bring new information that will surely help to distinguish the different stellar and planetary formation processes. This is the purpose of our work, in which we analyzed a sample of 200 stars to homogeneously measure the photospheric parameters and C abundances. ",
        "conclusions": "The preliminary results on C abundances are presented here for a large number of nearby solar-like stars, based on homogeneous photospheric parameters obtained from spectra with high signal-to-noise ratio and high resolution. Our analysis used public spectra from the ELODIE database, which represent about 90\\% of all data. The remaining 10\\% include many stars with detected planets and having many observations that shall be analyzed in the same way as soon as they become available to the scientific community, since they will contribute to more reliable conclusions."
    },
    "0910/0910.0327_arXiv.txt": {
        "abstract": "We investigate structures of hybrid stars, which feature quark core surrounded by a hadronic matter mantle, with  super-strong toroidal magnetic fields in full general relativity. Modeling the equation of state (EOS) with a first order transition by bridging the MIT bag model for the description of quark matter and the nuclear EOS by Shen et al., we numerically construct thousands of the equilibrium configurations to study the effects of the phase transition. It is found that the appearance of the quark phase can affect distributions of the magnetic fields inside the hybrid stars, making the maximum field strength about up to $30$ \\% larger than for the normal neutron  stars. Using the equilibrium configurations, we explore the possible evolutionary paths to the formation of hybrid stars due to the spin-down of magnetized rotating neutron stars. We find that the energy release by the phase transition to the hybrid stars is quite large~($\\la 10^{52}~\\rm erg$) even for super strongly magnetized compact stars. Our results suggest that the strong gravitational-wave emission and the sudden spin-up signature could be observable signals of the QCD phase transition, possibly for a source out to Megaparsec distances. ",
        "introduction": "A very hot issue in hadronic and nuclear physics is to search the phase transition from baryons to their constitutes - deconfined quarks. Heavy ion colliders such as RHIC (Brookhaven) and LHC (CERN) are now on line to explore the QCD phase diagram for the high temperature and small baryon density regimes, for which lattice QCD calculations predict a smooth crossover to the QCD phase transition (see \\citet{stephanov04} for review).  Conversely for the low temperature and high baryon density regimes, compact stars are expected to provide a unique window on the phase transition at their extreme density with super-strong magnetic field. It has been suggested long ago that quark matter may exist in the interior of compact objects \\citep{itoh70,bodmer71,witten84}. Hybrid stars (and strange quark stars) are considered to be such objects, which feature quark cores surrounded by a hadronic matter mantle (or quark cores only) (for reviews, e.g., \\citet{weber,glendenning01}). Even if the relevant conditions could be reached in a laboratory in the near future \\citep{senger}, the conditions prevailing in the compact stars are different from those produced in accelerators, i.e., the matter is long-lived, charge neutral and in $\\beta$-equilibrium with respect to weak interactions. It is therefore important to investigate the properties of such ``exotic'' stars, providing hints about the main features of matter at those extreme conditions.  The formation of the quark cores in compact stars is expected to take place by a first-order phase transition \\citep{glen92,glendenning01}. Albeit still very uncertain (e.g., \\citet{horvath07}), such a transition would proceed by the conversion of initially metastable hadronic matter in the core into the new deconfined quark phase. The metastable phase could be formed as the central density of neutron stars exceeds a critical value, due to mass accretion, spin-down or cooling. Possible astrophysical cites are in the protoneutron stars during the collapse of of supernova cores (near the epoch of bounce \\citep{takahara,yasutake07} or at the late postbounce phase \\citep{gentile,nakazato08,sagert08}) and in old neutron stars accreting from their companions \\citep{benve,chau,lin06}. In whichever cases, the sudden nucleation of the exotic phase in the hadronic star will be accompanied by a core-quake and huge energy release of the gravitational binding energy. Such energy release has been proposed to explain the central engines of the gamma-ray bursts \\citep{bombaci,bere}. Possible observables of these transient phenomenon should be glitches, magnetar flares, and superbursts \\citep{alford}. In addition, the detection of gravitational waves \\citep{ioka03,yasutake07,lin06,abdikamalov08} and neutrinos \\citep{nakazato08,sagert08} generated at the moment of the phase transition, should supply us implications to unveil the mechanism of the phase transition.   Here it should be noted that most of these calculations/estimations concerning the phase transition inside compact stars are limited to a non-rotating case, in which the Tolman-Oppenheimer-Volkoff equation is solved to obtain their structures \\citep{haensel86,zdunik87,muto,drago04,yasutake09}. Exceptions are for \\citet{gourgoulhon99,yasutake05, zdunik07}, in which relativistic equilibrium configurations of rotating strange stars, beyond the so-called slow rotation approximations (see references in \\citet{glendenning01}), are constructed for estimating the energy release. In \\citet{gourgoulhon99,yasutake05}, hadron matter is assumed to be converted fully to quark matter, leading to the formation of strange quark stars like in \\citet{alcock86,benvenuto89,olesen91,lugones94}. However, it is recently pointed out that the strange quark stars could be ruled out by their too fast spin-down rates via gravitational radiations from r-modes instability \\citep{madsen}. The quasi-periodic oscillations (QPOs) of strange quark stars could not be reconciled with the observations \\citep{anna}. Moreover, \\citet{mallick09} has recently claimed that magnetars cannot convert to purely quark stars, but only to hybrid stars. These suggest us to pay attention to hybrid stars rather than stars made of purely deconfined quarks. In a very recent work by \\citet{zdunik08}, the phase transition of rotating hybrid stars is discussed, however, the equation of state (EOS) is made very idealistic, assumed to model a strong phase transition, seemingly less sophisticated than the EOSs in the recent literature cited above. In addition to rotation,  neutron stars observed in nature are magnetized with the typical magnetic field strength $\\sim 10^{11}$--$10^{13}$ G \\citep{Lyne}. The field strength is often much larger than the canonical value as $\\sim 10^{15}$ G for a special class of the neutron stars such as magnetars \\citep{Lattimer:2007,woods}. In a series of our recent papers \\citep{kiuchi08a,kiuchi08b,kiuchi08d}, we have studied equilibrium configurations of relativistic magnetized compact stars to apply for the understanding of their formations and evolutions, but with hadronic EOSs.  Above situations motivate us to investigate the structures of relativistic hybrid stars with magnetic fields, and by using them, to estimate the energy release at the phase transition. This study is posed as an extension to the study by \\citet{yasutake05}, in which the phase transition to the rotating strange stars was investigated. Due to the unavailability of the method to construct a fully general relativistic star with arbitrarily magnetic structures (namely both with toroidal and poloidal fields), we here consider the equilibrium with purely toroidal fields as in \\citet{kiuchi08b}. It is noted that the outcomes of the recent stellar evolution calculations \\citep{Heger:2004qp} and the MHD(magnetohydrodynamics) simulations of core-collapse supernovae (\\citet{kotake04,takiwaki04,obergaulinger06,luc2007,sawai2008,kiuchi08c,takiwaki09}), suggesting much dominance of the toroidal fields, are not in contradiction with the assumption. Following the scenario proposed in \\citet{yasutake05}, we consider the possible evolutionary tracks of a rapidly rotating and magnetized neutron star to a slowing rotating hybrid star due to the spin-down via gravitational radiation and/or magnetic braking. During the evolutions, the baryon mass and the magnetic field strength are taken to be constant for simplicity. The energy release can be estimated from the difference in the mass-energies between the hadronic star and the hybrid star along each sequence. By constructing thousands of equilibrium configurations, we hope to clarify the possible maximum energy release at the moment of the transition and discuss their astrophysical implications.  The paper is organized as follows. The method for constructing the EOS with the phase transition and the numerical scheme for the stellar equilibrium configurations are briefly summarized in Sec.~\\ref{sec2}.  Sec.~\\ref{sec:result} is devoted to showing numerical results. Summary and discussion follow in Sec.~\\ref{sec:summary}. In this paper, we use geometrical units with $G=c=1$. ",
        "conclusions": "\\label{sec:summary} In this study, we investigated structures of general relativistic hybrid stars containing super-strong magnetic fields. Pushed by the results of recent stellar evolution calculations and the outcomes of recent MHD simulations of core-collapse supernovae, we treated the toroidal fields only. Using the bulk Gibbs construction, we modeled the EOS with a first order transition by bridging the MIT bag model for the description of quark matter and the nuclear EOS by Shen et al. We found that the presence of the quark phase can affect the distribution of the magnetic fields inside the hybrid stars, leading to the enhancement of the field strength about $30$ \\% than for the normal neutron stars. Using the equilibrium configurations, we explored the possible evolutionary paths to the formation of hybrid stars due to the spin-down of magnetized and rotating neutron stars. For simplicity,  the total baryon mass and magnetic flux are taken to be conserved during the evolution but that the angular momenta are lost gravitational waves and/or magnetic breaking. We found that the energy release by the conversion to hybrid stars is typically $\\la 10^{52}~\\rm erg $, smaller than previously estimated for the conversion to the strange quark stars.  In association with the vast energy release, it is natural to expect the emissions of gravitational waves. Amplitude of the gravitational waves associated with the phase transition, $h$, is given as follows \\citep{pacheco98}, \\begin{eqnarray} h \\sim 2 \\left( \\frac{G \\alpha E_{\\rm tot}}{c^3 \\tau} \\right)^{1/2} \\frac{1}{r \\omega}. \\end{eqnarray} Here, $E_{\\rm tot}$, $\\alpha$, $\\tau$, $r$, and $\\omega$ are energy release, fraction of the energy release emitted in form of the gravitational waves, damping time scale of the gravitational wave, distance to the source, and angular velocity, respectively. To set an absolute upper bound, we could choose $ E_{\\rm tot} \\sim 10^{52}$ erg and  $\\omega \\sim 100$ rad s$^{-1}$ according to Table~\\ref{tab:Mb-const-200}. We assume that the source is in our galactic center of $r \\sim 10$ kpc and that the damping timescale and $\\alpha$ is $\\sim 100$ ms and $10^{-4}$, respectively inferred from \\citet{abdikamalov08}. The resulting amplitudes become as high as $h \\sim 10^{-18 \\sim -19}$ peaking at $\\sim$ kHz, which are surely within the detection limits for the laser interferometers on line such as LIGO, VIRGO, GEO600, and TAMA300. Considering the optimal sensitivity of the interferometers of $h \\sim 10^{-21}$ at $\\sim$ 1 KHz, such signals, far stronger than the ones from canonical core-collapse supernovae  (e.g., \\citet{kotake09a,kotake09b,murphy}), are possibly visible for a source out to Megaparsec distances. In combination with the signatures of the spin-up discussed before, we speculate that the waveforms of such strong gravitational waves could provide us hints of properties of QCD phase transition. To find the accurate waveforms, general relativistic simulations are indispensable (Kiuchi et al. in preparation). Moreover accurate predictions for neutrino emissions \\citep{gentile,nakazato08,sagert08} and their detectability are remained to be studied (e.g., \\citet{kawagoe}). This paper should help to construct the initial conditions for such studies.  \\citet{alcock86} claimed that the conversion from the hadronic to the mixed phase occurs instantaneously in the weak interaction timescales~(such as $ d+u \\leftrightarrow  s+u$) when the central density exceeds the critical value of $\\rho_1$ in Figures 1 and 2. If so, the newly-born neutron stars may convert to the hybrid stars in the postbounce phase, before they evolve from A to B illustrated in Figure \\ref{fig:senario}. In this case, the released energy may be used to power original core-collapse supernovae or more energetic supernovae such as hypernovae. Although the formation paths to hybrid stars are different from the ones discussed in this paper, the estimation method of the energy release here are also applicable if we implement the EOS at finite temperature and high lepton fraction \\citep{yasutake09}, which we plan to study as a sequel to this paper.  It has been pointed out that the super-strong magnetic field more than $10^{18}$~G can affect the stiffness of the EOS \\citep{broderick00}. As mentioned, the maximum magnetic field increases about $\\sim 30 \\%$ by the phase transition, while it increases only unit percent if the compact stars evolve to the non-rotating ones in absence of the transition. This suggests that the magnetic effects should be seriously taken into account to our models treating the phase transition. Especially for the quark phase, \\citet{fukushima07, noronha07} have recently found that the energy gaps of magnetic color-flavor-locked phase are oscillating functions of the magnetic field. Such effects are remained to be studied.  We should comment about the very high values of the magnetic fields, larger than $10^{18}$G. \\citet{ruderman00} suggested that such high magnetic fields can be brought to the stellar surface by buoyancy forces. According to their result, we estimate the buoyancy time scale in our models. The buoyancy force can be expressed as \\begin{eqnarray} F_b \\sim \\frac{B_\\phi^2}{8\\pi c_s^2} g, \\label{eq:fb} \\end{eqnarray} where $c_s$ and $g$ are the sound speed and the gravitational acceleration, respectively. Putting typical values inside compact objects into the equation, we roughly estimate the time scale $\\tau$ as \\begin{eqnarray} \\tau \\sim 10^{-3}-10^{-4} {\\rm s} \\left(\\frac{\\rho}{10^{15} {\\rm g/cm^3}}\\right)^{1/2} \\left(\\frac{R}{10^6 {\\rm cm}}\\right)^{1/2} \\left(\\frac{c_s}{10^{10}{\\rm cm/s}}\\right) \\left(\\frac{B_\\phi}{10^{18} {\\rm G}}\\right)^{-1} \\left(\\frac{g}{10^{14} {\\rm cm /s^{2}}}\\right)^{-1/2}, \\end{eqnarray} where $R$ is the typical size of compact stars. Clearly such magnetic fields are not stable against the buoyancy forces in stellar evolution time scale $\\sim 10^3$yr. Based on the MHD simulations in full general relativity, \\citet{kiuchi08c} recently showed that the toroidal configurations of neutron stars is really dynamically unstable due to the buoyant instability. But more important finding relevant to this study, is that the toroidal magnetic fields settle down to a new equilibrium state with the circular motions in the meridian plane. In the new equilibrium configurations, the toroidal fields almost equivalent to the strength before the onset of the instability, are expected to be much stronger that the poloidal ones. This suggests that our models presented here, albeit unstable to the buoyant instability, could be helpful to understand equilibrium configurations of magnetized hybrid stars. Such stability analysis is an important issue yet to be studied.    Our treatments of the phase transition should be more sophisticated. We plan to employ the so-called NJL model (e.g., \\citet{yasutake09}) instead of the simple MIT bag model. Moreover we shall take into account finite size effects \\citep{maruyama} such as quark-nuclear pasta structures. We considered cold compact objects, namely with zero-temperature and zero-neutrino fraction. We plan to extend this study to the finite temperature with the non-zero neutrino fraction, which should be useful for studying earlier evolutions of compact stars soon after their formation.  The evolution sequence we explored is a very simplified one. For more realistic estimations, one needs the two dimensional evolutionary calculation in full general relativistic framework, not an easy job, due to the treatment of convection, combustion, cooling and/or heating processes, nucleosynthesis on the surface, and so forth. Applying the method recently reported by \\cite{pons}, we think that we could study the magneto-thermal evolution of neutron stars to the hybrid stars. By doing so, we hope to investigate the peculiar properties in the light curves, which has been pointed out to be another observationally visible \\citep{page00,blaschke00,blaschke01,grigorian05}. All of such studies should be indispensable to pave the way for the understanding of the long-veiled phase-transition physics from the astrophysical phenomena."
    },
    "0910/0910.1117_arXiv.txt": {
        "abstract": "The diffusion of astrophysical magnetic fields in conducting fluids in the presence of turbulence depends on whether magnetic fields can change their topology via reconnection in highly conducting media. Recent progress in understanding fast magnetic reconnection in the presence of turbulence is reassuring that the magnetic field behavior in computer simulations and turbulent astrophysical environments is similar, as far as magnetic reconnection is concerned. This makes it meaningful to perform MHD simulations of turbulent flows in order to understand the diffusion of magnetic field in astrophysical environments. Our studies of magnetic field diffusion in turbulent medium reveal interesting new phenomena. First of all, our three-dimensional MHD simulations initiated with anti-correlating magnetic field and gaseous density exhibit at later times a de-correlation of the magnetic field and density, which corresponds well to the observations of the interstellar media. While earlier studies stressed the role of either ambipolar diffusion or time-dependent turbulent fluctuations for de-correlating magnetic field and density, we get the effect of {\\it permanent} de-correlation with one fluid code, i.e. without invoking ambipolar diffusion. In addition, in the presence of gravity and turbulence, our three-dimensional simulations show the decrease of the magnetic flux-to-mass ratio as the gaseous density at the center of the gravitational potential increases. We observe this effect both in the situations when we start with equilibrium distributions of gas and magnetic field and when we follow the evolution of collapsing dynamically unstable configurations. Thus the process of turbulent magnetic field removal should be applicable both to quasi-static subcritical molecular clouds and cores and violently collapsing supercritical entities. The increase of the gravitational potential as well as the magnetization of the gas increases the segregation of the mass and magnetic flux in the saturated final state of the simulations, supporting the notion that the reconnection-enabled diffusivity relaxes the magnetic field + gas system in the gravitational field to its minimal energy state. This effect is expected to play an important role in star formation, from its initial stages of concentrating interstellar gas to the final stages of the accretion to the forming protostar. In addition, we benchmark our codes by studying the heat transfer in magnetized compressible fluids and confirm the high rates of turbulent advection of heat obtained in an earlier study. ",
        "introduction": "Astrophysical flows are known to be turbulent and magnetized. The specific role played by MHD turbulence in different branches of astrophysics is still highly debated, but it is generally regarded as important. In particular, for the interstellar medium (ISM) and star formation, the role of turbulence has been discussed in many reviews (see \\citealt{elmegreen2004,mckee2007}). The opinion on the role of magnetic field in these environments vary from magnetic field being regarded as absolutely dominant in the processes (see \\citealt{tassis2005,galli2006}) to moderately important, as in super-Alfv\\'enic models of star formation (see \\citealt{padoan2004}).  The vital question that frequently permeates these debates is the diffusion of the magnetic field in astrophysical fluids. The conductivity of most of the astrophysical fluids is high enough to make the Ohmic diffusion negligible on the scales involved, which means that the ``frozen-in'' approximation is a good one for many astrophysical environments. However, without considering diffusive mechanisms that can violate the flux freezing, one faces problems attempting to explain many observational facts. For example, simple estimates show that if all the magnetic flux is brought together with the material that collapses to form a star in molecular clouds, then the magnetic field in a proto-star should be several orders of magnitude higher than the one observed in T-Tauri stars (this is the ``magnetic flux problem'', see \\citealt{galli2006} and references therein, for example).  To address the problem of the magnetic field diffusion both in the partially ionized ISM and in molecular clouds, researchers usually appeal to the ambipolar diffusion concept (see \\citealt{mestel1956,shu1983}). The idea of the ambipolar diffusion is very simple and may be easily exemplified in the case of gas collapsing to form a protostar. As the magnetic field is acting on charged particles only, it does not directly affect neutrals. Neutrals move under the gravitational pull but are scattered by collisions with ions and charged dust grains which are coupled with the magnetic field. The resulting flow dominated by the neutrals will be unable to drag the magnetic field lines and these will diffuse away through the infalling matter. This process of ambipolar diffusion becomes faster as the ionization ratio decreases and therefore, becomes more important in poorly ionized cloud cores.  \\citet{shu2006} have explored the accretion phase in low-mass star formation and concluded that there should exist  an effective diffusivity about 4 orders of magnitude larger than Ohmic diffusivity in order to an efficient magnetic flux transport to occur.  They have argued that ambipolar diffusion could work, but only under  rather special circumstances like, for instance, considering particular dust grain sizes. In other words, at the moment it is unclear if  ambipolar diffusion is  really high enough to solve the magnetic flux transport problem in collapsing flows.  Does magnetic field remain absolutely frozen-in within highly ionized astrophysical fluids? The answer to this question affects the description of numerous essential processes in the interstellar and intergalactic gas.  Magnetic reconnection was appealed in \\citet{lazarian2005} as a way of removing magnetic flux from gravitating clouds, e.g. from star-forming clouds. That work referred to the reconnection model of \\citet{lazarian1999} and \\citet{lazarian2004} for the justification of the concept of fast magnetic reconnection in the presence of turbulence. The advantage of the scheme proposed by \\citet{lazarian2005} was that robust removal of magnetic flux can be accomplished both in partially and fully ionized plasma, with only marginal dependence on the ionization state of the gas\\footnote{The rates were predicted to depend on the reconnection rate, which according to \\citet{lazarian2004} very weakly depends on the ionization degree of the gas.}. The concept of ``reconnection diffusion'' introduced in \\citet{lazarian2005} is relevant to our understanding of many basic astrophysical processes. In particular, it suggests that the classical textbook description of molecular clouds supported both by hourglass magnetic field and turbulence is not self-consistent. Indeed, turbulence is expected to induce ``reconnection diffusion'' which should enable fast magnetic field removal from the cloud.  However, in the absence of numerical confirmation of the fast reconnection, the scheme of magnetic flux removal through ``reconnection diffusion'' as opposed to ambipolar diffusion stayed somewhat speculative.  Fortunately, it has been recently shown numerically (see \\citealt{kowal2009}) that three-dimensional magnetic reconnection in turbulent fluid follows the predictions of the reconnection model of \\citet{lazarian1999} and therefore, is fast. This naturally increased the interest to the ``reconnection diffusion'' (see \\citealt{lazarian2009}). Motivated by this fact, here we perform simulations aiming to gain understanding of the diffusion of magnetic field induced by turbulence. As \\citet{kowal2009} tested magnetic reconnection in fully ionized gas in the present paper we shall focus our efforts on  one fluid MHD simulations.  We shall compare ``reconnection diffusion'' with the ambipolar diffusion and discuss the effects of ambipolar diffusion qualitatively focusing on the  comparison of our results on ``reconnection diffusion'' with those on ``ambipolar turbulent diffusion'' described in \\citet{heitsch2004}. The latter study  reported the enhancement of ambipolar diffusion in the presence of turbulence, which raised the question of how important is the simultaneous action of turbulence and ambipolar diffusion and whether turbulence alone, i.e. without any effect from ambipolar diffusion, can equally well induce de-correlation of magnetic field and density.  Our work on ``reconnection diffusion'' should  also be distinguished from the research on the de-correlation of magnetic field and density within compressible turbulent fluctuations. \\citet{cho_lazarian2002,cho_lazarian2003} performed three-dimensional MHD simulations and reported the existence of separate turbulent cascades of Alfv\\'en and fast modes in strongly driven turbulence as well as a cascade of slow modes driven by Alfv\\'enic cascade. Slow modes in magnetically dominated plasma are associated with density perturbations with marginal perturbation of magnetic fields, while the same is true for fast modes in weakly magnetized or high beta plasmas. Naturally, these two modes de-correlate magnetic fields and density on the crossing time of the wave. This was the effect studied in more detail in one-dimensional both analytically and numerically by \\citet{passot2003}, who stressed that the enhancements of magnetic field strength and density may correlate and anti-correlate in turbulent interstellar gas within the fluctuations and this can introduce the dispersion of the mass-to-flux ratios within the turbulent volume. Each of the fluctuations provide a {\\it transient} change of the pointwise magnetization. In the absence of other effects, e.g. related to the thermal instability, the de-correlation is reversible. In comparison, the ``reconnection diffusion'', similar to the ambipolar diffusion, deals with the {\\it permanent} de-correlation of magnetic field and density. Acting alone, the ``reconnection diffusion'' increases entropy making magnetic field-density de-correlation irreversible.   What are the laws that govern \\textit{magnetic field diffusion} in turbulent magnetized fluids? Could those affect our understanding of basic interstellar and star formation processes? These are the questions that we address in this paper.    In this study, we try to understand the diffusion of magnetic field in a couple of idealized models in the presence of turbulence. We explore setups both with and without gravity and compare the diffusion of magnetic field with that of a passive scalar. In the context of star formation an important issue that we will address is an alternative way of decreasing the magnetic flux-to-mass ratio without appealing to ambipolar diffusion. We claim that since turbulence in astrophysics is really ubiquitous, our results should be widely applicable. We also perform simulations including turbulent heat diffusion, which allow us to compare results obtained with our code with those in the literature.  This paper is organized as follows. In Section 2, we draw the theoretical grounds about  fast magnetic reconnection. In Section 3, we describe the numerical code employed.  In Section 4, we present the results concerning the diffusion of magnetic field in a setup without external gravitational forces. In Section 5, we present the results of our experiments of diffusion of magnetic field in the presence of a gravitational field. In Section 6, we discuss our results and compare with previous works. In Section 7, we discuss the accomplishments and limitations of our  present study. In Section 8, we discuss our findings in the context of strong turbulence theory, and finally in Section 9, we summarize our conclusions. While our work is focused on the diffusion of magnetic fields, we address in the Appendix the heat transport in magnetized turbulent plasma.  We confirm with higher resolution the results in \\citet{cho2003a} that the heat advection within turbulent flows in the presence of magnetic fields is very similar to that induced by hydrodynamic turbulence. ",
        "conclusions": "Through this work we have performed the comparison of our results with the study by \\citet{heitsch2004}. Below, we provide yet another outlook of the connection of that study with the present paper. We also discuss the work by \\citet{shu2006}, which was the initial motivation of our study of the diffusion of magnetic field in the presence of gravity.   \\subsection{Comparison with \\citet{heitsch2004}: Ambipolar Diffusion Versus Turbulence and 2.5-dimensional Versus Three-dimensional}  In view of the astrophysical implications, the comparison between our results and those of \\citet{heitsch2004} calls for the discussion on how ambipolar diffusion and turbulence interact to affect the magnetic field diffusivity. In particular, \\citet{heitsch2004} claim that a new process ``turbulent ambipolar diffusion'' (see also \\citealt{zweibel2002}) acts to induce fast magnetic diffusivity.  At the same time, our results do not seem to exhibit less magnetic diffusivity than those of \\citet{heitsch2004} in spite of the fact that we do not have ambipolar diffusion. How can this be understood? We propose the following explanation. In the absence of ambipolar diffusion, the turbulence propagates to smaller scales making small-scale interactions possible. On the other hand, ambipolar diffusion affects the turbulence, increasing the damping scale. As a result, the ambipolar diffusion acts in two ways, in one to increase the small-scale diffusivity of the magnetic field, in another is to decrease the turbulent small-scale diffusivity and these effects essentially compensate each other\\footnote{A possible point of confusion is related to the difference of the physical scales involved. If one associates the scale of the reconnection with the thickness of the Sweet--Parker layer, then, indeed, the ambipolar diffusion scale is much larger and therefore the reconnection scale gets irrelevant. However, within the LV99 model of reconnection, the scale of reconnection is associated with the scale of magnetic field wandering. The corresponding scale depends on the turbulent velocity and is not small.}.  In other words, if we approximate the turbulent diffusivity by $(1/3) V_{\\text{inj}} L_{\\text{inj}}$, where $V_{\\text{inj}}$ and $L_{\\text{inj}}$ are the injection velocity and the injection scale for strong MHD turbulence (see LV99, \\citealt{lazarian2006}), respectively, the ambipolar diffusivity acting on small scales will not play any role and the diffusivity will be purely ``turbulent''. If, however, the ambipolar diffusion coefficient is larger than $V_{\\text{inj}} L_{\\text{inj}}$, then the Reynolds number of the steered flow may become small for strong MHD turbulence to exist and the diffusion is purely ambipolar in this case. We might speculate that this leaves little, if any, parameter space for the ``turbulent ambipolar diffusion'' when turbulence and ambipolar diffusion synergetically enhance diffusivity, acting in unison. This point should be tested by three-dimensional two-fluid simulations exhibiting both ambipolar diffusion and turbulence.  In view of our findings one may ask whether it is surprising to observe the ``reconnection diffusion'' of magnetic field being independent of ambipolar diffusion. We can appeal to the fact well known in hydrodynamics, namely, that in a turbulent fluid the diffusion of a passive contaminant does not depend on the microscopic diffusivity. In the case of high microscopic diffusivity, the turbulence provides mixing down to a scale $l_1$ at which the microscopic diffusivity both, suppresses the cascade and ensures efficient diffusivity of the contaminant. In the case of low microscopic diffusivity, turbulent mixing happens down to a scale $l_2\\ll l_1$, which ensures that even low microscopic diffusivity is sufficient to provide efficient diffusion. In both cases the total effective diffusivity of the contaminant is turbulent, i.e. is given by the product of the turbulent injection scale and the turbulent velocity. This analogy is not directly applicable to ambipolar diffusion, as this  is a special type of diffusion and magnetic fields are different from passive contaminants. However, we believe that our results show that to some extent the concept of turbulent diffusion developed in hydrodynamics carries over (due to fast reconnection) to magnetized fluid.  \\subsection{Transient De-correlation of Density and Magnetic Field}  Magneto-sonic waves are known to create transient changes of the density and magnetic field correlation. In the case of turbulence the situation is less clear, but the research in the field suggests that the decomposition of the turbulent motions into basic MHD modes is meaningful and justified even for high amplitude motions (\\citealt{cho2003a}). Thus the claim in \\citet{passot2003} that even in the limit of ideal MHD, turbulence can {\\it transiently} affect the magnetic field and density correlations is justified. However, the process discussed in this paper is different in the sense that the de-correlation we describe here is {\\it permanent} and it will not disappear if the turbulence dissipates. In a sense, as we showed above, ``reconnection diffusion'' is similar to the ambipolar and Ohmic diffusion. It is a dissipative diffusion process, which does require non-zero resistivity, although this resistivity can be infinitesimally small for the LV99 model of fast reconnection in the presence of turbulence (see Section 2).  \\subsection{Relation to \\citet{shu2006}: Fast Removal of Magnetic Flux During Star Formation}  As discussed in \\citet{shu2006}, the sufficiency of the ambipolar diffusion efficiency for explaining observational data of accreting proto-stars is questionable. At the same time, they found that the required dissipation is about 4 orders of magnitude larger than the expected Ohmic dissipation. Thus they appealed to the hyper-resistivity concept in order to explain the higher dissipation of magnetic field.  We feel, however, that the hyper-resistivity idea is poorly justified (see criticism of it in \\citealt{lazarian2004} and \\citealt{kowal2009}). At the same time, fast three-dimensional ``reconnection diffusion'' can provide the magnetic diffusivity that is adequate for fast removing of the magnetic flux. This is what, in fact, was demonstrated in the present set of numerical simulations.  It is worth mentioning that, unlike the actual Ohmic diffusivity, ``reconnection diffusion'' does not transfer the magnetic energy directly into heat. The lion share of the energy is being released in the form of kinetic energy, driving turbulence (see LV99). If the system is initially laminar, this potentially result in flares of reconnection and the corresponding diffusivity. This is in agreement with LV99 scheme where a more intensive turbulence should induce more intensive turbulent energy injection and lead to the unstable feeding of the energy of the deformed magnetic field. However, the discussion of this effect is beyond the scope of the present paper.  Similar to \\citet{shu2006}, we expect to observe the heating of the media. Indeed, although we do not expect to have Ohmic heating, the kinetic energy released due to magnetic reconnection is dissipated locally and therefore we expect to observe heating in the medium. Our setup for gravity can be seen as a toy model representing the situation in \\citet{shu2006}. In the broad sense, our work confirms that a process of magnetic field diffusion that does not rely on ambipolar diffusion is efficient.  We accept that our setup assuming an axial gravitational field is a very simple and ignores complications that could arise from using a nearly spherical potential of the self-gravitating cloud. The periodic boundary conditions give super-stability to the system, and do not allow inflow (or outflow) of material/magnetic field as we expect in a more realistic accretion process. However, our experiments can give us qualitative insights. They show that the turbulent diffusion of the magnetic field can remove magnetic flux from the central region, leading to a lower flux-to-mass ratio in regions of higher gravity compared with that of lower gravity.  We chose parameters to the simulations such that the system is not initially unstable to the Parker--Rayleigh--Taylor (PRT) instability. Although the PRT instability could be present in real accretion systems and could help to remove magnetic field from the core of gravitational systems, its presence would make the interpretation of the results more difficult and we wanted to analyze only the turbulence role in the removal of magnetic flux. However, it is possible that this instability had been also acting due to local changes of parameters due to the turbulent motion. To ensure that the transport of magnetic flux is being caused by injection of turbulence only, we stopped the injection after a few time steps in some experiments and left the system to evolve. When we did this, the changes in the profile of the magnetic field and the other quantities stopped.  We showed that the higher the strength of the gravitational force, the lower the flux-to-mass ratio is  in the central region (compared with the mean value in the computational domain). This could be understood in terms of the potential energy of the system. When the gravitational potential well is deeper, more energetically  favorable is the pile up of matter near the center of gravity, reducing the total potential energy of the system. When the turbulence is increased, there is an initial trend to remove more magnetic flux from the center (and consequently more inflow of matter into the center), but for the highest value of the turbulent velocity in our experiments, there is a trend to remove material (together with magnetic flux) from the center, reducing the role of the gravity, due to the fact that the gravitational energy became small compared to the kinetic energy of the system. Our results also showed that when the gas is less magnetized (higher $\\beta$, or higher values of the Alfv\\'enic Mach number $M_{A}$), the reconnection diffusion of magnetic flux is more effective, but the central flux-to-mass ratio relative to external regions is  smaller for more magnetized models (low $\\beta$), compared to less magnetized models. That is, the contrast $B / \\rho$ between the inner and outer radius is higher for lower $\\beta$ (or $M_{A}$).  If the turbulent diffusivity of magnetic field may explain the results in \\citet{shu2006}, one may wonder whether one can remove magnetic field by this way not only from the class of systems studied by \\citet{shu2006}, but also from less dense systems. For instance, it is frequently assumed that only ambipolar diffusion is important for the evolution of subcritical magnetized clouds \\citep{tassis2005}. Our study indicates that this conclusion may be altered in the presence of turbulence.  This point, however, requires further  careful study, which is beyond the scope of the present paper. In the future, we intend to study a more realistic model, e.g., with open boundary conditions and more realistic gravitational potentials.       Motivated by a vital problem of the dynamics of magnetic fields in astrophysical fluids, in particular, by the magnetic flux removal in star formation,  in this paper we have numerically studied the diffusion of magnetic field both in the absence and in the presence of gravitational potential. Recent work on validating the idea of LV99 model of reconnection supports our assertion that our results obtained at moderate resolution represent the dynamics of turbulent magnetic field lines in astrophysics. Our findings obtained on the basis of three-dimensional MHD numerical simulations can be briefly summarized as follows:  1. In the absence of gravitational potential the ``reconnection diffusion'' removes strong anti-correlations of magnetic field and density that we impose at the start of our simulations. The system after several turbulent eddy turnover times relaxes to a state with no clear correlation between magnetic field and density, reminiscent of the observations of the diffuse ISM by \\citet{troland1986}.  2. Our simulations that started with a quasi-static equilibrium in the presence of a gravitational potential, revealed that the turbulent diffusivity induces gas to concentrate at the center of the gravitational potential, while the magnetic field is efficiently pushed to the periphery. Thus the effect of the magnetic flux removal from collapsing clouds and cores, which is usually attributed to ambipolar diffusion effect, can be  successfully accomplished without ambipolar diffusion, but in the presence of turbulence.  3. Our simulations that started in a state of dynamical collapse induced by an external gravitational potential showed that in the absence of turbulence, the flux-to-mass ratio is preserved for the collapsing gas. On the other hand, in the presence of turbulence, fast removal of magnetic field from the center of the gravitational potential occurs. This may  explain the low magnetic flux-to-mass ratio observed in stars compared to the corresponding ratio of the interstellar gas.  4. As an enhanced Ohmic resistivity to remove magnetic flux from cores and accretion disks has been appealed in the literature, e.g., by \\citet{shu2006}, we have also compared models with a turbulent fluid and models without turbulence but with substantially enhanced Ohmic diffusivity. We have shown that, in terms of the magnetic flux removal, the reconnection diffusion can mimic the effect of an enhanced Ohmic resistivity.  5. In addition, our results extend earlier findings of \\citet{cho2003a} for heat advection by magnetized turbulence. We show that heat advection can be parameterized by the product of the turbulence injection scale and the turbulent velocity for a range of Alfv\\'enic and sonic Mach numbers."
    },
    "0910/0910.3865_arXiv.txt": {
        "abstract": "{% The first supersoft source (SSS) identification with an optical nova in \\m31 was based on ROSAT observations. Twenty additional X-ray counterparts (mostly identified as SSS by their hardness ratios) were detected using archival ROSAT, \\xmm\\  and \\chandra\\  observations obtained before July 2002. Based on these results optical novae seem to constitute the major class of SSS in \\m31. An analysis of archival \\chandra\\  HRC-I and ACIS-I observations obtained from July 2004 to February 2005 demonstrated that \\m31  nova SSS states lasted from months to about 10 years. Several novae showed short X-ray outbursts starting within 50 d after the optical outburst and lasting only two to three months. The fraction of novae detected in soft X-rays within a year after the optical outburst was more than 30\\%. Ongoing optical nova monitoring programs, optical spectral follow-up and an up-to-date nova catalogue are essential for the X-ray work. Re-analysis of archival nova data to improve positions and find additional nova candidates are urgently needed for secure recurrent nova identifications. Dedicated \\xmm/\\chandra\\  monitoring programs for X-ray emission from optical novae covering the center area of \\m31  continue to provide interesting new results (e.g. coherent 1105 s pulsations in the SSS counterpart of nova M31N 2007-12b). The SSS light curves of novae allow us -- together with optical information -- to estimate the mass of the white dwarf, of the ejecta and the burned mass in the outburst. Observations of the central area of \\m31\\ allow us -- in contrast to observations in the Galaxy -- to monitor many novae simultaneously and proved to be prone to find many interesting SSS and nova types.} ",
        "introduction": "The outbursts of classical novae (CNe) are caused by explosive hydrogen burning on the white dwarf (WD) surface of a cataclysmic variable, a close binary system with transfer of material from a main sequence star to the WD. After about $10^{-7}-5\\times10^{-4}$~M$_{\\odot}$ of H-rich material are transferred to the WD, ignition under degenerate conditions takes place in the accreted envelope and a thermonuclear runaway is initiated \\citep[see e.g.][]{1998ApJ...494..680J,2005ApJ...623..398Y}. As a consequence, the envelope expands and causes the brightness of the star to increase to maximum luminosities of up to $\\sim 10^{5}$~L$_{\\odot}$. A fraction of the envelope is ejected, while a part of it remains in steady nuclear burning on the WD surface. This powers a supersoft X-ray source (SSS) which can be observed as soon as the expanding ejected envelope becomes optically thin to soft X-rays, with the spectrum of a hot ($T_{eff}: 10^{5}-10^{6}$~K) WD atmosphere \\citep[][]{1991ApJ...373L..51M}. The duration of the SSS phase is inversely related to the WD mass while the time of appearance of the SSS is determined by the mass ejected in the outburst and the ejection velocity. Models of the post-outburst WD envelope show that steady H-burning can only occur for envelope masses smaller than $\\sim 10^{-5}$~M$_{\\odot}$ \\citep[][]{2005A&A...439.1061S,1998ApJ...503..381T}, and the observed evolution of the SSS in V1974 Cyg has been successfully modeled by an envelope of $\\sim 2\\times10^{-6}$~M$_{\\odot}$ \\citep[][]{2005A&A...439.1057S}. WD envelope models show that the duration of the SSS state also depends on the metalicity of the envelope, so the monitoring of the SSS states of CNe provides constraints also on the chemical composition of the post-outburst envelope. \\citet[][]{2006ApJS..167...59H} have developed envelope and wind models that simulate the optical and X-ray light curves for several WD masses and chemical compositions.  Accreting WDs in recurrent novae (RNe) are good candidates for type Ia supernovae (SNe) as RNe are believed to contain massive WDs. However, one of the main drawbacks to make RNe convincing progenitors of SNe-Ia was their low fraction in optical surveys \\citep[][]{1994ApJ...423L..31D}. In the case of CNe the ejection of material in the outburst makes it difficult to follow the long-term evolution of the WD mass. For some CNe there is a disagreement between theory and observations regarding the ejected masses, with observational determinations of the mass in the ejected shell larger than predicted by models. The duration of the SSS phase provides the only direct indicator of the post-outburst envelope mass remaining on the WD in RNe and CNe. In the case of CNe with massive WDs, the SSS state is very short ($<$100 d) and could have been easily missed in previous surveys. CNe with short SSS state are additional good candidates for SNe-Ia progenitors which makes determining their frequency very important.   Nevertheless, the number and duration of SSS states observed in optical novae are small: in a systematic search of X-ray emission from CNe in the ROSAT archive \\citep[][]{2001A&A...373..542O}, found only three novae with SSS emission in X-rays from a total of 39 CNe observed less than ten years after the outburst with SSS phases lasting between 400~d and 9~yr. The \\chandra\\ and \\xmm\\ observatories have detected SSS emission for several more novae; but only a limited number of observations have been performed for each source, providing little constraints on the duration of the SSS state. Of specific interest was the monitoring of the recurrent nova RS~Oph in spring 2006 with the Swift satellite which clearly determined the end of the SSS state less than 100~d after outburst \\citep[see e.g.][]{2006ATel..838....1O} which suggests a WD mass of 1.35 M$_{\\odot}$ \\citep[][]{2006ApJ...651L.141H}. The observations of the Galactic nova V458 Vul \\citep[detected as highly variable SSS $\\sim$400 d after outburst, see e.g.][]{2008ATel.1721....1D} and nova V2491 Cyg \\citep[starting its SSS state 36~d after the optical outburst, see e.g.][]{2008ATel.1573....1N} demonstrated that each Galactic nova seems to have its own peculiarities.  The small number of novae found to exhibit a SSS state, and the diversity of the duration of this state (from 10 years down to few weeks) present one of the big mysteries in the study of hydrogen burning objects over the last years. Despite an extensive target of opportunity program with \\chandra\\ and \\xmm\\ (of order 3 dozen observations during the last 4 years), little progress has been made in constraining the duration of SSS states, or to even putting constraints on the long term evolution of accreting WDs in binary systems. This now may change with the monitoring campaign for Galactic novae with the \\swift\\ satellite (see Osborne et al. in this issue).  \\begin{figure} \\includegraphics[width=82mm,clip=]{fovima_2008.eps} \\caption{DSS1 image of M\\,31: the SuperLOTIS and Skinakas field of view (FoV) are overlaid as big squares. Circles show the \\chandra\\ HRC~I and \\xmm\\ EPIC PN fields. Novae with outbursts after 2000 are indicated by green dots. Big circles indicate X-ray detected novae (extracted from our \\m31\\ nova web page, see http://www.mpe.mpg.de/$\\sim$m31novae/opt/m31/M31\\_table.html).} \\label{fovima} \\end{figure} ",
        "conclusions": ""
    },
    "0910/0910.1098_arXiv.txt": {
        "abstract": "We analyse recently acquired near-infrared {\\em Hubble} Space Telescope imaging of the GOODS-South field to search for star forming galaxies at $z\\approx 7.0$. By comparing WFC\\,3 $0.98\\mu$m $Y$-band images with ACS $z$-band ($0.85\\mu$m) images, we identify objects with colours consistent with Lyman break galaxies at z$\\simeq$6.4-7.4. This new data covers an area five times larger than that previously reported in the WFC3 imaging of the {\\em Hubble} Ultra Deep Field, and affords a valuable constraint on the bright end of the luminosity function. Using additional imaging of the region in the ACS $B$, $V$ and $i$-bands from GOODS v2.0 and the WFC3 $J$-band we attempt to remove any low-redshift interlopers. Our selection criteria yields 6 candidates brighter than $Y_{\\rm AB}=27.0$, of which all except one are detected in the ACS $z$-band imaging and are thus unlikely to be transients. Assuming all 6 candidates are at $z\\approx 7$ this implies a surface density of objects brighter than $Y_{\\rm AB}=27.0$ of 0.30$\\pm$0.12\\,arcmin$^{-2}$, a value significantly smaller than the prediction from $z\\approx 6$ luminosity function. This suggests continued evolution of the bright end of the luminosity function between $z=6\\to 7$, with number densities lower at higher redshift. ",
        "introduction": "In recent years our understanding of the high-redshift galaxy population has rapidly expanded with the discovery of star-forming galaxies within the first billion years ($z>5$), through the Lyman break technique using broad-band imaging (e.g. Stanway, Bunker \\& McMahon 2003; Dickinson et al.\\ 2004) and searches for Lyman-$\\alpha$ emission with narrow-band filters (e.g., Ouchi et al.\\ 2008, Ota et al.\\ 2008), and most recently gamma-ray bursts (e.g. Tanvir et al.\\ 2009, Salvaterra et al. 2009). At $z\\sim 6$, the brightest ($z<26.5$) Lyman-break ``$i$-band drop-out\" galaxies have been confirmed spectroscopically (e.g. Bunker et al.\\ 2003; Stanway et al.\\ 2004ab; Dow-Hygelund et al.\\ 2007; Vanzella et al.\\ 2009) through their Lyman-$\\alpha$ emission, confirming the validity of this photometric redshift selection.  In recent weeks, the installation of the new WFC3 camera on the {\\em Hubble} Space Telescope ({\\em HST}), which has an infrared channel with a  large field of view, has enabled the Lyman break technique to be pushed to $z\\sim 7-10$, revealing for the first time significant numbers of galaxy candidates close to the reionization epoch (Bunker et al.\\ 2009; Bouwens et al.\\ 2009; Oesch et al.\\ 2009; McLure et al.\\ 2009; Yan et al.\\ 2009). However, the first data to be released was the extremely deep single pointing on the Hubble Ultra Deep Field ($\\approx 4\\,{\\rm arcmin}^{2}$), reaching objects as faint at mag$_{AB}=28.5$ ($6\\,\\sigma$) in $Y$-, $J$- and $H$-bands. To probe the rare galaxies at the bright end of the luminosity function at $z\\sim 7-8$ requires a larger field of view. There have been some shallower wider area searches using ground based observations (e.g. Hickey et al. 2009, Ouchi et al. 2009, Castellano et al. 2009), but the depths probed are typically $Y_{AB}<26$ (equivalent to $L_{UV}>2\\,L^*$) and the numbers of robust candidates are small. Increasing the number of $z\\approx 7$ candidates over a wide range of magnitudes is critical to exploring the shape of the luminosity function. This is particularly important as there is strong evidence for evolution of luminosity function, with suggestions that at high redshift there is a larger relative contribution to the stellar mass and UV luminosity density from sub-$L_{*}$ systems. At $z\\sim 7$ we are close to the epoch of reionization, and an open question is the mechanism by which reionization is achieved; if we are to address the contribution of the UV from star-forming galaxies, then quantifying the luminosity function is vital (along with the escape fraction of ionizing photons from galaxies, and the hardness of the UV spectral slope).  In this paper we present first results from {\\em HST}/WFC3 imaging of some of the GOODS-South field, covering an area 5 times that of the WFC3 images of the {\\em Hubble Ultra Deep Field}, and reaching $Y_{AB}=27.0$ ($6\\,\\sigma$), probing luminosities around $L^*_{UV z\\sim 6}$. In conjunction with the GOODS v2.0 Advanced Camera for Surveys (ACS) images with $B,V,i,z$ bands (Giavalisco et al.\\ 2004) we search for objects which are much brighter in the $Y$-band WFC3 filter at $1\\,\\mu$m than the $z$-band $0.9\\,\\mu$m, and are undetected at shorter wavelength. These ``$z$-drops'' are candidate $z\\sim 7$ Lyman break galaxies, and the greater area of this new dataset (compared with the Hubble Ultra Deep Field WFC3 images) means that we are likely to find brighter objects more amenable to future spectroscopic confirmation.  The structure of this paper is as follows. In Section 2 the imaging data and the construction of catalogues  is described. In Section 3 we describe our candidate selection and discuss the observed  surface density of $z\\approx 7$ galaxies, comparing with the expected number from a range of luminosity functions. Our conclusions are presented in Section 4. Throughout we adopt the standard ``concordance'' cosmology of $\\Omega_{M}=0.3$, $\\Omega_{\\Lambda}=0.7$, and use $h_{70}=H_{0}/70 {\\rm kms^{-1}\\,Mpc^{-1}}$. All magnitudes are on the AB system (Oke \\& Gunn 1983). ",
        "conclusions": "In this work we have searched for star-forming galaxies at $z\\approx 7$ utilising the Lyman-break technique on newly acquired F098M $Y$-band images from WFC3 on the {\\em Hubble} Space Telescope. Through the comparison of these images to existing {\\em Hubble} ACS F850LP $z$-band images we identified objects with red colours, $(z-Y)>0.8$, indicative of a break in the spectrum. We explore an area five times larger than the recent WFC\\,3 imaging of the {\\em Hubble} Ultra Deep Field. The new wider-field data in GOODS-South probes down to $Y_{AB}=27$, equivalent to $\\approx L^*_{UV}$ at $z=7$.  Using additional imaging (ACS $B$, $B$, $i$-bands from GOODS v2.0, and WFC3 F125W $J$-band) we removed contaminating objects which were either detected in the bluer ACS bands or with observed near-infrared colours inconsistent with high-redshift star forming galaxies.  This selection criteria left 6 candidates down to a limiting magnitude $Y_{AB}<27.0$ (equivalent to a star formation rate of $4.5\\,{\\rm M_{\\odot}\\,yr^{-1}}$ at $z=7$) of which all but one were detected in the $z$-band (and are thus likely not to be transients). This implies a surface density of objects brighter than $Y_{\\rm AB}=27.0$ of 0.30$\\pm$0.12$\\, {\\rm arcmin}^{2}$; a value smaller than both the prediction based on the observed $z\\approx 3$ and $z\\approx 6$ luminosity functions, suggesting continued evolution of the LF beyond $z=6$.  Knowledge of the surface density of dropouts in this magnitude range is crucial in constraining the luminosity function, as current estimates of the $z\\approx 7$ luminosity function indicate that an $L_{*}$ galaxy has a magnitude of $Y\\sim 27$ (e.g. Bouwens et al. 2008). Determining the UV luminosity function is crucial in order to address whether star-forming galaxies could plausibly have provided the Lyman continuum photons necessary to reionize the Universe.  Given the difficulty that ground-based surveys face in probing faintward of $Y=26$, larger-area surveys with {\\em HST}/WFC3 ($\\sim 100$ arcmin$^2$) probing to similar depths ($Y(6\\sigma)\\sim 27.2$) as the dataset presented in this paper will deliver a factor of 2 improvement in the poisson uncertainty and a significant decrease in the uncertainty due to cosmic variance.  \\subsection*{Acknowledgements}  We would like to thank the annoyamous referee for their timely response and useful suggestions. Additionally we thank Mark Lacy for comments given during the construction of the manuscript, and Richard McMahon, Jim Dunlop, Ross McLure, Masami Ouchi, Bahram Mobasher and Michele Cirasuolo for many useful discussions about Lyman break galaxies at high redshift. Based on observations made with the NASA/ESA Hubble Space Telescope, obtained from the Data Archive at the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS 5-26555. These observations are associated with programme \\#GO/DD-11359. We are grateful to the  WFC\\,3 Science Oversight Committee for making their Early Release Science observations public. SW, DS \\& ES acknowledges funding from the U.K.\\ Science and Technology Facilities Council. MJJ acknowledges the support of a RCUK fellowship. SL is supported by the Marie Curie Initial Training Network ELIXIR of the European Commission under contract PITN-GA-2008-214227."
    },
    "0910/0910.4168_arXiv.txt": {
        "abstract": "We apply the axisymmetric orbit superposition modeling to estimate the mass of the supermassive black hole and dark matter halo profile of NGC~4649. We have included data sets from the Hubble Space Telescope, stellar, and globular cluster observations. Our modeling gives $\\mbh= (4.5\\pm 1.0) \\times 10^9\\Msun$ and $\\mlvobs=8.7 \\pm 1.0$ (or $\\ml_V=8.0\\pm 0.9$ after foreground Galactic extinction is corrected). We confirm the presence of a dark matter halo, but the stellar mass dominates inside the effective radius. The parameters of the dark halo are less constrained due to the sparse globular cluster data at large radii. We find that in NGC~4649 the dynamical mass profile from our modeling is consistently larger than that derived from the X-ray data over most of the radial range by roughly 60\\% to 80\\%. It implies that either some forms of non-thermal pressure need to be included, the assumed hydrostatic equilibrium may not be a good approximation in the X-ray modelings of NGC~4649, or our assumptions used in the dynamical models are biased. Our new $\\mbh$ is about two times larger than the previous published value; the earlier model did not adequately sample the orbits required to match the large tangential anisotropy in the galaxy center. If we assume that there is no dark matter, the results on the black hole mass and $\\mlvobs$ do not change significantly, which we attribute to the inclusion of {\\it HST} spectra, the sparse globular cluster kinematics, and a diffuse dark matter halo. Without the {\\it HST} data, the significance of the black hole detection is greatly reduced. ",
        "introduction": "\\label{sec:intro}   Most nearby galaxies harbor supermassive black holes at their centers. Correlations between black hole (BH) mass and host galaxy properties \\citep{mag_etal_98,geb_etal_00,fer_mer_00,har_rix_04} have been used extensively in theoretical models in order to understand growth of the black hole and galaxy \\citep[e.g.,][]{hop_etal_08}.  The latest work from Hopkins et al. suggest that a main role of a black hole is to halt star formation in the galaxy when the black hole is large enough, thereby causing the dichotomy in colors \\citep[e.g.,][]{bell_08}. While there is still an active debate as to the relative role of AGN feedback versus star formation feedback, there is a consensus that physical mechanisms for black hole growth are very important for understanding mass growth in galaxies. A concern is that the black hole correlations may have significant systematic biases, both from kinematics with poor spatial resolution and models that do not adequately include the full mass profile (see discussion in \\citealt{gul_etal_09}).  Dynamical modeling of galaxies using orbit superposition offers one of the best estimates on the black hole mass \\citep[e.g.,][]{rix_etal_97,van_etal_98,cre_etal_99,geb_etal_00,geb_etal_03,val_etal_04,tho_etal_04,tho_etal_05,sio_etal_09}. Assuming axisymmetry, these models do not limit the form of the allowed velocity anisotropies. Thus, the stellar orbital structure resulting from the dynamical modeling provides a unique window on the mass growth process in the massive systems, as long as the model assumptions are valid.  A particular systematic bias is shown in \\citet{geb_tho_09} where they find that the black hole mass can be underestimated in the most massive galaxies if the dark halo is not included. They find a degeneracy between the dark halo and black hole mass, since without the dark halo the stellar mass-to-light ratio is overestimated which subsequently decreases the required contribution of the black hole to constrain the central kinematics. In M87, \\citet{geb_tho_09} find that the black hole mass goes from $2.5\\times10^9$ to $6.4\\times10^9~\\Msun$ by simply running models including a dark halo. Furthermore, the uncertainties do not overlap, implying a large systematic effect. In M87, however, this degeneracy is strong since the black hole is not well resolved by the kinematic data. For galaxies with well spatially-resolved kinematics, we do not expect the degeneracy to be as significant.  In addition to studying the black hole mass, there is a strong need to study the shapes of dark matter profiles. There is still little consensus for those measured in individual galaxies (e.g., PNe from \\citealt{rom_etal_03}, stellar light from \\citealt{for_geb_08}, X-rays from \\citealt{gas_etal_07,chu_etal_08,hum_etal_08,hum_etal_09}, and globular clusters from \\citealt{bri_etal_06,hwa_etal_08}). The impressive work using gravitational lensing to measure the average dark matter profiles (e.g., \\citealt{man_etal_06a,man_etal_06b}) has been able to reach out to nearly 1~Mpc. However, these results need to be compared to measurements based on individual galaxies. It is important to understand the galaxy-to-galaxy scatter in the profiles and whether there are environmental effects.  This paper is part of an extensive campaign to measure both the black hole mass and the dark matter profile simultaneously, to examine the possible bias of the dark matter profile on the black hole mass estimate for galaxies with various profiles, and to compare with the gravitational potentials derived with other independent techniques such as the X-rays and weak lensing studies. Initially, we focus on the more massive galaxies, but it will be important to extend future analysis over a large mass range. \\citet{geb_tho_09} report results for M87, the most massive nearby elliptical, and in this paper we focus on NGC~4649. Both galaxies are giant ellipticals with a central surface brightness ``core''.  In this paper we present the axisymmetric orbit superposition models for NGC~4649 (M60), combining data from the Hubble Space Telescope ({\\it HST}), stellar, and globular cluster observations.  NGC~4649 is a giant elliptical with low surface brightness located in a subclump to the east of the main Virgo concentration \\citep{for_etal_04}.  NGC~4649 has been studied extensively for its total mass profile in recent X-ray modelings \\citep[e.g.,][]{hum_etal_06,gas_etal_07,hum_etal_08}, and globular cluster studies \\citep[e.g.,][]{bri_etal_06}. The goals of our study are to place NGC~4649's black hole mass estimate on a more solid footing, to infer the properties of its dark matter halo, and more importantly to offer an independent cross-check using different dynamical tracers to the previous studies on NGC~4649.  We assume a distance to NGC~4649 of 15.7 Mpc. At this distance, 1$\\arcsec$ corresponds to 76$\\;$pc. ",
        "conclusions": "\\label{sec:conclusions}  We model the dynamical structure of NGC~4649 using the high resolution data sets from {\\it HST}, stellar, and globular cluster observations. Our main new results are:  1. Our modeling gives $\\mbh= 4.5 \\pm 1.0 \\times 10^9\\Msun$ and $\\mlvobs=8.7 \\pm 1.0$.  Our new $\\mbh$ of NGC~4649 is about a factor of 2 larger than the previous result. We find that the earlier model did not adequately sample the orbits required to match the large tangential anisotropy in the galaxy center..  2. We confirm the presence of a dark matter halo in NGC~4649, but the stellar mass dominates inside the effective radius. The parameters of the dark halo especially the core radius are less constrained due to the sparse globular cluster data at large radii.  3. Unlike in the case of M87, the black hole mass from the dynamical modeling is not biased as much by the inclusion of a dark matter halo, because high-resolution {\\it HST} spectra are available for NGC~4649, the globular cluster kinematics are sparse, and the halo is not as dominant inside the effective radius $R_e$ as that of M87.  4. We find that in NGC~4649 the dynamical mass profile from our modeling is consistently larger than that derived from the X-ray data over most of the radial range by about 70\\%. It implies that either some forms of non-thermal pressure need to be included, the assumed hydrostatic equilibrium may not be a good approximation in the X-ray modeling, or the assumptions used in our dynamical modeling create a bias."
    },
    "0910/0910.5956.txt": {
        "abstract": "A study of metal enrichment of the intergalactic medium (IGM) using a series of smooth particle hydrodynamics (SPH) simulations is presented, employing models for metal cooling and the turbulent diffusion of metals and thermal energy. An adiabatic feedback mechanism was adopted where gas cooling was prevented on the timescale of supernova bubble expansion to generate galactic winds without explicit wind particles.  The simulations produced a cosmic star formation history (SFH) that is broadly consistent with observations until z $\\sim$ 0.5, and a steady evolution of the universal neutral hydrogen fraction ($\\Omega_{\\rm H I}$) that compares reasonably well with observations. The evolution of the mass and metallicities in stars and various gas phases was investigated.  At z=0, about 40\\% of the baryons are in the warm-hot intergalactic medium (WHIM), but most metals (80\\%-90\\%) are locked in stars.  At higher redshifts the proportion of metals in the IGM is higher due to more efficient loss from galaxies.  The results also indicate that IGM metals primarily reside in the WHIM throughout cosmic history, which differs from simulations with hydrodynamically decoupled explicit winds. The metallicity of the WHIM lies between 0.01 and 0.1 solar with a slight decrease at lower redshifts.  The metallicity evolution of the gas inside galaxies are broadly consistent with observations, but the diffuse IGM is under enriched at  z $\\sim$ 2.5.  Galactic winds most efficiently enrich the IGM for halos in the intermediate mass range $10^{10}$M$_{\\sun}$ - $10^{11}$ M$_{\\sun}$. At the low mass end gas is prevented from accreting onto halos and has very low metallicities. At the high mass end,  the fraction of halo baryons escaped as winds declines along with the decline of stellar mass fraction of the galaxies.  This is likely because of the decrease in star formation activity and decrease in wind escape efficiency. Metals enhance cooling which allows WHIM gas to cool onto galaxies and increases star formation.  Metal diffusion allows winds to mix prior to escape, decreasing the IGM metal content in favour of gas within galactic halos and star forming gas.  Diffusion significantly increases the amount of gas with low metallicities and changes the density-metallicity relation. ",
        "introduction": "The intergalactic medium (IGM) contains most of the baryons in the Universe and it provides the fuel for galaxies to form stars in which metals are produced. In turn, supernovae and galactic winds enrich the IGM with metals,  while stars and  active galactic nuclei (AGN) emit UV photons.  This interplay between the IGM and galaxies, mediated by metal cooling in the presence of UV, regulates the formation of stars in the universe.  The evolution and enrichment history of the IGM provides a record of this interplay.  Observations of metal absorption lines (e.g., C III , C IV, Si III, S IV and O VI) in quasar spectra show that the intergalactic medium (IGM) far outside large galaxies ($\\rho/\\rho_{mean} < 10$) is enriched \\citep[e.g.][]{Songaila96, Dave98, Ellison00, Schaye00, Pettini03, Schaye03, Aguirre04, Simcoe04}. There is evidence for enrichment extending back to $z > 5$ \\citep{Pettini03, Simcoe06} though metallicities do not evolve much from z=4  to z=2 \\citep{Schaye03}.  Since metals are created in stars inside galaxies, those in the IGM must have escaped from galaxies. Exactly how this occurs is unclear. It typically assumed that galactic winds, driven by star formation, dominate IGM enrichment \\citep{Aguirre07}. Winds can be either launched as the ejecta of a large number of co-existing supernova (SN) explosions \\citep{Heckman90}, driven by the injection of momentum by SN and stellar winds or by radiation pressure from starbursts and AGN \\citep{Murray05}. Observations of galactic winds \\citep[e.g.,][]{Heckman01, Pettini01} have found that galactic wind velocities range from hundreds to thousands of km/s and the mass loss rates are comparable to star formation rates.  Wind material has a complex, multiphase structure but most metals are expected to be entrained in the hot phase \\citep[][and references therein]{Veilleux05}.  Although detailed hydrodynamical simulations of the interstellar medium (ISM) in galaxies \\citep[e.g,,][]{MacLow99, Strickland00, Williams02} have been able to generate galactic outflows and have explored various properties of winds and the gas dynamics in different phases, current cosmological simulations lack the resolution to launch or track winds directly. Hot, low density SN bubbles are unresolved in such simulations which initially led to an overcooling problem that produced unrealistically concentrated simulated galaxies \\citep{Navarro97}. As a result, various ``subgrid'' stellar feedback and wind models have emerged. These models serve two functions: to regulate star formation and the properties of the ISM and to redistribute gas (and newly formed metals) both within and into the environment around galaxies.  There are three main approaches, energetic feedback, kinetic feedback and modifications to the effective equation of state which behaves similarly to an increased effective pressure.  Energetic feedback in its simplest form involves simply adding the stellar feedback as thermal energy, but this suffers from overcooling \\citep{Katz96}.  Kinetic feedback \\citep[e.g.,][]{Navarro93, SH03, Oppenheimer06, DallaVecchia08} converts part of the SN energy into kinetic energy in the gas. The effectiveness of kinetic feedback is strongly dependent on the resolution and hydrodynamic method.  \\citet{SH03} argued that regulated star formation creates an effective pressure in the ISM and this was modelled directly in the \\textsc{gadget} code as part of a recipe for regulated star formation.  This approach leads to a strongly hydrodynamically-coupled, multiphase ISM that does not naturally produce galactic outflows.  To combat this the authors added a ``superwind'' model where fluid elements in the star forming region are ejected at fixed speed and are also hydrodynamically decoupled until they leaves the galaxy.  \\citet{Oppenheimer06} modified the model in a manner referred to as the momentum-driven wind scenario so that the velocity of the wind and the mass loading factor were related to the velocity dispersion of the host galaxy.  The \\textsc{gadget} code with superwind feedback prescriptions has been widely used in various problems such as damped Lyman-$\\alpha$ (DLA) absorbers \\citep{Nagamine04a, Nagamine04b} and the enrichment of the IGM at high and low redshifts \\citep{Oppenheimer06, Oppenheimer09}.  According to these works, superwind feedback is essential to suppress overproduction of stars in galaxies and to reproduce the cosmic SFH at high redshift. It also increases the local fraction of the warm-hot intergalactic medium (WHIM) to a sufficient percentage (40\\% to 50\\%) to account for the ``missing baryons'' at z = 0. Although aspects of the model compare well with observations, some components do not.  For example, the feedback may eject a large amount of cool gas from the galactic disks, which results in a low neutral hydrogen mass density $\\Omega_{\\rm H I}$ at z $<$ 2 \\citep{Nagamine04a}. Also, the interaction between winds and the ISM is usually not modeled in these simulations.   \\citet{DallaVecchia08} found that the ISM plays an important role in regulating the amount of wind that escapes and the morphology of the galaxies. In their model winds are not hydrodynamically decoupled, which naturally allows for variable mass loading.  A refined version of energetic feedback is adiabatic feedback which treats the overcooling problem by inhibiting gas from cooling until the hot SN bubbles can be resolved \\citep[e.g.,][]{Thacker00,   Kay02, SommerLarsen03, Stinson06}.  This is the approach used in this work. The pressure of the hot gas accelerates the ISM to generate winds.  In the high resolution limit, this method approaches direct ISM modeling.  Though energetic feedback is often referred to as supernova feedback, it can be used to model several types of stellar feedback such as winds and locally deposited radiation energy.  The essential quantity is the energy injection rate as a function of the mass in stars and the current age of the stellar population.  \\citet{Theuns02} used the \\citet{Kay02} adiabatic feedback model that turned off cooling for 10 Myrs for the feedback gas in their cosmological simulations. They found enough metals were carried by strong winds to produce C IV absorption lines that agreed with observations.    \\citet{Aguirre05} used the same simulation to compare the optical depth of C IV and C III absorption lines from simulations with observations.  The properties of the enriched IGM are not only affected by winds but also by gas cooling.  Winds can enrich galactic halos and the IGM so that metal cooling significantly increases the cooling rates. \\citet{Aguirre05} found that their simulated metal enriched gas was too hot ($10^{5} \\sim 10^{7}$ K) and suggested that a lack of metal cooling was responsible for discrepancies between simulated and observed C IV absorption.  \\citet{Oppenheimer06} found better agreement when they included the \\citet{Sutherland93} metal cooling model.  The same model was used by \\citet{Choi09}, who investigated the effect of metal cooling on galaxy growth and found that it increases the local star forming efficiency and enhances accretion onto galaxies.  However, \\citet{Sutherland93} did not include photoionization due to a ultraviolet (UV) radiation background which strongly affects the ionization states of metal species and changes the cooling rates. This is investigated in detail in the current work and also in \\citet{Wiersma09a}.  Another important aspect of metal enrichment is the mixing of metals between the wind and the surrounding gas.  The interstellar medium is highly turbulent and SN explosions are likely to be a major driver of the turbulence \\citep{MacLow04}.  In addition, large velocity shear (such as between a wind and a gaseous halo) naturally generates turbulence and mixing. Turbulent mixing redistributes metals and thermal energy between the wind fluid and the ambient gas.  This changes the metallicity, temperature and future evolution of the gas. While metal mixing is expected in strong outflows, it is still unclear how mixing impact the IGM. For example, observations by \\citet{Schaye07} found compact ($\\sim$ 100 pc), transient C IV absorbers that are highly enriched, suggesting poor chemical mixing at small scales.  These absorbers were interpreted as enriched clumpy medium embedded within hot galactic wind fluids. If velocity shear is the major mechanism for turbulent mixing between winds and the surroundings,  then this poor mixing could be explained if the clouds are carried by hot winds at the same speed. However to investigate this in detail, one must resolve wind structures, which is beyond current cosmological simulations. In this work we will focus on subgrid turbulent mixing models in cosmological context.   In SPH simulations (which represent the majority of work in this area), the fluid is modeled by discrete particles. This implies that newly injected metals are locked into specific particles.  For example, it was found that the distribution of metals from SPH simulations is too inhomogeneous compared with observations \\citep{Aguirre05}. To assess the potential importance of mixing, \\citet{Wiersma09} used SPH-smoothed metallicities and compared it with conventional particle metallicities, and found smoothing is able to generate significantly more material with low metallicities.  This approach cannot capture the spread of metals over time with its impact on cooling and the thermal history of the gas. Directly modeling the turbulent ISM within a cosmological simulation is far beyond current capabilities.  We employed a variant of the \\citet{Smagorinsky63} subgrid turbulent diffusion model, in which unresolved turbulent mixing is treated as a shear-dependent diffusion term. Metal cooling was calculated based on the diffused metals so that its non-local effects could be investigated.  In this work, we present an analysis of a series of SPH cosmological simulations that incorporated adiabatic stellar feedback, detailed metal cooling and turbulent mixing to study the evolution and enrichment of the IGM. The feedback model was kept simple, following the adiabatic stellar feedback approach of \\citet{Stinson06}. This model has been calibrated via numerous galaxy formation studies \\citep[e.g., ][]{Governato07}. No additional wind prescriptions were used.  With this approach, outflows arise from stellar feedback within the ISM and there is no distinction between the feedback that regulates star formation and that which drives galactic outflows.  Thus this work establishes a baseline for the effectiveness of moderate stellar feedback coupled with key physical process absent from other work to reproduce the properties of the IGM.  These results may be compared with explicit wind models. A further goal of this paper is to separately quantify the impact of metal cooling and turbulent mixing on the SFH, the global properties and the evolution of the IGM and its enrichment.  In this first paper of a series, we present general results. The properties of specific metal absorbers in simulated quasar spectra will be presented in a second paper.  This paper is organized as follows.  Section~\\ref{method} describes the models for cosmological hydrodynamics, star formation, supernova feedback, metal cooling and metal diffusion.  Section~\\ref{global} examines the cosmic SFH, global H I fraction and Ly-$\\alpha$ decrement in order to calibrate our models. Section~\\ref{mass_metal_evo} focuses on the evolution of the baryonic mass, metal fractions and metallicities in stars and different gas phases.  We compare those results to the observed metal fractions and metallicities at different epochs, and with the simulations using different subgrid feedback models from \\citet{Oppenheimer06}, \\citet{Dave07} and \\citet{Wiersma09}.  In section~\\ref{phasediagram} we analyze the distribution of mass and metallicity in the density-temperature phase diagram at z = 0.  In section~\\ref{wind} we characterize our wind efficiency as a function of galaxy mass to obtain a better understanding of how different phases of the IGM get enriched.  Where relevant,  we have included detailed analysis of the effects of metal cooling and diffusion, and comparisons with observations.  In the final section~\\ref{summary}, we summarize and discuss the broader implications. ",
        "conclusions": "\\label{summary}  We investigated the enrichment of the intergalactic medium with SPH cosmological simulations using an adiabatic stellar feedback model. The simulations incorporated a self-consistent metal cooling model with an ultraviolet (UV) ionizing background along with metal diffusion that models the turbulent mixing in the IGM and the ISM.  It was found that the UV background significantly alters the metal cooling rates at all temperatures from 100 K to $10^{9}$ K.  Above $10^{4}$ K it decreases the cooling rate and shifts the cooling peak to higher temperature, while below $10^{4}$ K the UV increases the metal cooling rates due to the increase of free electrons.  The simulations produced an SFH broadly consistent with observations to redshift z $\\sim$ 0.5,  and a steady cosmic total neutral hydrogen fraction ($\\Omega_{\\rm H I}$) that compare relatively well with observations, although possible discrepancies in the observed $\\Omega_{\\rm H I}$ at z $\\sim$ 2-3 allow for more vigorous mass-loss.  This demonstrates that adiabatic feedback can moderate SF while maintaining a regular supply of H I. The evolution of the mean flux decrement in the Ly-$\\alpha$ forest in our simulations is consistent with observations to z $\\sim$ 3-4, if the magnitude of the UV background is lowered with respect to the \\cite{HM05} rates in \\textsc{cloudy}.  As the universe evolves, there is a rapid increase in the amount of warm-hot intergalactic medium (WHIM) and a decrease in the cooler diffuse IGM.  At z=0, about 40\\% of the mass is in WHIM, consistent previous simulations with different methods \\citep{Cen06b, Oppenheimer06}. The metal content of the Universe evolves from the largest fraction being in the IGM to the majority residing in star forming gas, to ultimately being locked in stars at the present day. These trends reflect more effective wind escape at high redshift.  IGM metals primarily reside in the WHIM, unlike \\citet{Oppenheimer06} and \\citet{Dave07}, whose metals largely reside in cool gas, as the wind gas starts cool with kinetic feedback and less likely to be heated up when the wind material is decoupled from the ISM.  Our result is however in agreement with \\citet{Wiersma09} who also used kinetic feedback but with non-decoupled wind models. %Compared to the observations, our simulations produce a similar %estimated metal budget at z=0.  At z $\\sim$ 2-3, our results are consistent %with observations for stellar and IGM (diffuse plus WHIM) metal fractions, %but have significantly more metals in the star forming gas. The mean metallicities of stars, star forming gas, galactic halo gas and the cold diffuse IGM all increase with time, but those of the WHIM and the ICM remain mostly constant with a slight decreasing trend. The metallicity of the WHIM stays between 0.01 to 0.1 $Z_{\\sun}$. The metallicity of the ICM is similar to the WHIM, which is smaller than the observed value, 0.2-0.5 $Z_{\\sun}$ \\citep{Aguirre07}, possibly due to absent AGN feedback.   The metallicity evolution of the gas in galaxies compares well with DLA and sub-DLA observations from \\citet{Prochaska03} and the metallicity of the diffuse IGM at z $\\sim$ 2.5 is less than the observations from \\citet{Schaye03} and \\citet{Aguirre08}, suggesting higher resolution, or additional mechanisms for enriching the diffuse IGM are probably necessary.  We characterized our galactic mass-loss and wind generation.  For the current adiabatic feedback model, winds are most efficient for galaxies in the intermediate mass range of $10^{10}$M$_{\\sun}$ to about $10^{11}$M$_{\\sun}$.  Below $10^{10}$M$_{\\sun}$ gas is likely to be prevented from accreting due to UV heating and remains as low-metallicity gas.  Above about $10^{11}$ M$_{\\sun}$ fraction of wind gas decreases, possibly because the wind escape efficiency decreases with increasing halo potentials, or because the decline of star formation activities especially for massive halos. Most winds were hot when generated, but the ones expelled from intermediate mass range galaxies, having temperatures $\\sim 10^{5} - 10^{6}$K, could cool through metal lines and become diffuse IGM rather than WHIM.  We investigated the effect of metal cooling and diffusion on the SFH, the evolution of $\\Omega_{\\rm H I}$ and the evolution of mass and metals in different phases.  For metal diffusion, we further studied its effect in the density-temperature phase diagram at z =0 and the distribution of baryons as function of the galaxy mass.   Metals significantly enhance the cooling of the WHIM, allowing the gas to cool and join galactic disks.  Metals also enable cooling below $10^{4}$K.  Metal cooling decreases the mass and metal fractions of the WHIM while increasing the metals in stars, halos and SF gas and increasing the SFR by 20\\% and $\\Omega_{\\rm H I}$ by 17\\% at z =0. With realistic diffusion included, metals mix between winds and surrounding gas before they leave the galaxies, decreasing the metal content in the WHIM and diffuse IGM but increasing it in the galactic halo and star forming gas. It prevents enriched, hot winds from creating highly-enriched low density regions, and makes the density-metallicity relation smoother so it follows a nearly log-linear relation in the density range $log(\\rho/\\rho_{mean}) = 0-4$.  We performed two additional simulations with one of which having 8 times more particles,  to investigate the convergence of our results and the resolution dependence of our model.  Full convergence is not seen in most of our results, however the results from the high resolution run in SFH, $\\Omega_{\\rm H I}$ evolution, Ly$\\alpha$ decrement evolution are still broadly consistent with observational data, and the differences caused by resolution decrease with redshift. Also, the basic results of metal fraction evolution trend and enrichment efficiency in different gas phases remain unchanged. The impact of resolution can be seen in three aspects. Firstly, with high resolution the amount of neutral gas (H I) and SFR increase, hence the stellar mass fraction and the WHIM mass fraction related to wind activities.  The difference is largest at high redshift. Secondly, the metal fraction locked up in stars increases with resolution while decreasing that in the ISM and IGM in general, suggesting a decline in wind efficiency. Thirdly, with high resolution the diffuse IGM and halos have a higher metal fraction at the cost of the WHIM metals.  In future work, we will make a more detailed comparison with metal absorption line observations.  We will also examine the detailed behaviour of various wind models and the effect on the immediate environment of galaxies."
    },
    "0910/0910.5114_arXiv.txt": {
        "abstract": "In this article we identify and discuss various statistical and systematic effects influencing the astrometric accuracy achievable with MICADO, the near-infrared imaging camera proposed for the 42-metre European Extremely Large Telescope (E-ELT). These effects are instrumental (e.g. geometric distortion), atmospheric (e.g. chromatic differential refraction), and astronomical (reference source selection). We find that there are several phenomena having impact on $\\sim$100$\\mu$as scales, meaning they can be substantially larger than the theoretical statistical astrometric accuracy of an optical/NIR 42m-telescope. Depending on type, these effects need to be controlled via dedicated instrumental design properties or via dedicated calibration procedures. We conclude that if this is done properly, astrometric accuracies of 40$\\mu$as or better -- with 40$\\mu$as/yr in proper motions corresponding to $\\approx$20~km/s at 100~kpc distance -- can be achieved in one epoch of actual observations. ",
        "introduction": "The future optical/near-infrared European Extremely Large Telescope (E-ELT; see, e.g., Gilmozzi \\& Spyromilio \\cite{gilmozzi2008}), which is designed with a 42-metre aperture, will offer a substantial improvement in angular resolution compared to existing facilities. At wavelengths $\\lambda=2\\mu$m, diffraction-limited resolutions of  $\\Theta\\simeq10$mas will be achieved. In terms of angular resolution in the near infrared, the E-ELT will outperform existing 8--10m-class telescopes like the VLT or Keck by factors of $\\approx$4--5 and the future James Webb Space Telescope (JWST) by factors of $\\approx$7. This increase in angular resolution should translate into a corresponding improvement in astrometric accuracy.  In order to exploit the E-ELT's resolution, a German-Dutch-Italian-French consortium\\footnote{The MICADO collaboration includes: MPE Garching, Germany; USM Munich, Germany; MPIA Heidelberg, Germany; NOVA (a collaboration of the universities of Leiden, Groningen, and ASTRON Dwingeloo), The Netherlands; OAPD Padova (INAF), Italy; LESIA Paris, France} proposed the Multi-AO Imaging Camera for Deep Observations (MICADO) in February 2008. As the spatial resolution of any ground-based observatory is initially limited by the atmospheric seeing, MICADO will be equipped with a multi-conjugate adaptive optics (MCAO) system for achieving the diffraction limit of the 42m-telescope. This system uses three natural and six laser guide stars for correcting the atmospheric turbulence in a wide ($>2'$) field of view  (Diolaiti et al. \\cite{diolaiti2008}). Images will be recorded by an array of 4$\\times$4 near-infrared (NIR) HAWAII-4RG detectors with 4096$\\times$4096 pixels each, covering a FOV of 53'. The instrument is sensitive to the wavelength range $0.8-2.5\\mu$m, thus covering the I, Y, J, H, K bands. For astrometric experiments the use of the data analysis software {\\sl Astro-WISE} (Valentijn et al. \\cite{valentijn2007}) is foreseen.  In order to achieve its science goals (see Sect.~2 for details), MICADO needs to reach a stable (time scales of years) astrometric accuracy of approximately 50$\\mu$as. At present 8--10m class telescopes, accuracies of $\\approx$0.5\\% of a resolution element can be reached regularly (e.g. Fritz et al. \\cite{fritz2009b}). Therefore from simple scaling of results our goal \\emph{a priori} appears reasonable. However, at levels of the order of 100$\\mu$as there are several sources of statistical and systematic errors which need to be taken into account carefully. In this article we discuss those effects and analyse strategies to bypass them. We conclude that reaching an astrometric accuracy of better than 50$\\mu$as is highly challenging in terms of instrument design and data calibration but feasible.  MICADO's astrometric performance should be of the same magnitude as that of the future astrometry space mission GAIA (e.g. Jordan \\cite{jordan2008}). GAIA will achieve accuracies better than $\\approx$50$\\mu$as only for bright (V$<$15.5) targets and only at the end of its mission. MICADO is expected to achieve this accuracy for targets with $K_{\\rm AB}<26$. Other space missions like SIM PlanetQuest (e.g. Edberg et al. \\cite{edberg2007}) or JASMINE (e.g. Gouda et al. \\cite{gouda2007}) also aim specifically at bright targets in order to reach accuracies of $\\approx$10$\\mu$as (at best).  For illustration purposes, Fig.~\\ref{fig_eelt_gc} shows simulated observations of the nuclear star cluster of the Milky Way using both present day 8-10m class telescopes and E-ELT/MICADO. Physical parameters of the star cluster (stellar density profile, luminosity function) are taken from Genzel et al. \\cite{genzel2003}. We discuss technical details of our simulations in Sect. 4.1. These maps demonstrate the impressive progress to be expected with MICADO.   \\begin{figure} \\includegraphics[height=8.8cm]{f01a.eps} \\includegraphics[height=8.8cm]{f01b.eps} \\caption{An illustration of the expected performance of the E-ELT/MICADO system. These simulated maps show the central $1''\\times1''$ (i.e. 8000AU$\\times$8000AU) of the nuclear star cluster of the Milky Way at 2.2$\\mu$m. \\emph{Top panel}: The target region as observed with present day 8-10m class telescopes. The diffraction-limited resolution is $\\approx$50mas. For comparison with actual observations, see, e.g., Genzel et al. (2003), Ghez et al. (2005). \\emph{Bottom panel}: The same field as seen by MICADO. The angular resolution is $\\approx$10mas. The improvement in detail and depth is obvious.} \\label{fig_eelt_gc} \\end{figure}   Although this study is set up for the specific case of MICADO, most of its results are valid in general and therefore of interest beyond the E-ELT community.  This paper is organised as follows. In Section~2, we discuss the science cases identified for MICADO. In Section~3, we review the concepts and techniques of accurate astrometry. In Section~4, we identify and analyse sources of systematic errors one by one and describe methods for minimizing those errors. We provide a summary of our results and an overall error budget in Section~5 and present our conclusions in Section~6. ",
        "conclusions": "In this article we have studied the capabilities expected for the NIR imager MICADO for the future 42-m European Extremely Large Telescope with respect to accurate astrometry. A variety of science cases requires long-term astrometric accuracies of $\\approx$50$\\mu$as. We discuss and quantify ten effects that potentially limit the astrometric accuracy of MICADO. We conclude that the systematic accuracy limit for astrometric observations with MICADO is $\\sigma_{\\rm sys}\\approx40\\mu$as. We find that astrometry at this accuracy level with MICADO requires the fulfillment of several conditions:  \\begin{itemize}  \\item  All images, regardless of their distance in time, need to be combined via full coordinate transforms of second order or higher.  \\item  MICADO needs to be equipped with an astrometric calibration mask for monitoring the instrumental distortion. The pixel scale of the camera should not exceed the 3~mas/pix used in the current design.  \\item  Astrometric observations require decent integration times of at least 30 minutes per epoch. This is unavoidable in order to average out atmospheric tilt jitter. When using high-z galaxies as astrometric reference points, integration times up to about 10 hours can be necessary.  \\end{itemize}   It is noteworthy that the effects discussed in this article already affect observations collected with present 8m-class telescopes. In his exhaustive analysis of NIR images obtained with VLT/NACO, Fritz \\cite{fritz2009} has been able to detect signatures of chromatic differential refraction and differential tilt jitter in his astrometric dataset. He concludes that taking into account these effects can improve the accuracies down to few hundred $\\mu$as. This agrees with the findings of Lazorenko \\cite{lazorenko2006} and Lazorenko et al. \\cite{lazorenko2007} who analyze seeing-limited optical R-band ($\\lambda_{\\rm center}=655$nm) images taken with VLT/FORS1+2. They conclude that they are able to achieve astrometric precisions (but not accuracies) of $\\approx$100$\\mu$as by using a special scheme for scheduling observations and dedicated coordinate transforms (although they neglect instrumental geometric distortion).  The analysis we provide here is set up for the specific case of E-ELT/MICADO, but parts of our results are valid in general. This study should thus contain valuable information for other future 30--40m telescopes. As some of the effects we discuss are actually observed in present day 8m-class telescope data, our analysis might also be helpful for the calibration of data already taken. We therefore expect that our work is of interest well beyond the E-ELT community."
    },
    "0910/0910.3397_arXiv.txt": {
        "abstract": "We compare the abundances of various chemical species as derived with 3D hydrodynamical and classical 1D stellar atmosphere codes in a late-type giant characterized by \\teff$=3640$\\,K, $\\log g=1.0$, $\\MoH=0.0$. For this particular set of atmospheric parameters the 3D--1D abundance differences are generally small for neutral atoms and molecules but they may reach up to 0.3--0.4\\,dex in case of ions. The 3D--1D differences generally become increasingly more negative at higher excitation potentials and are typically largest in the optical wavelength range. Their sign can be both positive and negative, and depends on the excitation potential and wavelength of a given spectral line. While our results  obtained with this particular late-type giant model suggest that 1D stellar atmosphere models may be safe to use with neutral atoms and molecules, care should be taken if they are exploited with ions. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.4262_arXiv.txt": {
        "abstract": "We use photometric and spectroscopic observations of the eclipsing binary V69-47~Tuc to derive the masses, radii, and luminosities of the component stars. Based on measured systemic velocity, distance, and proper motion, the system is a member of the globular cluster 47~Tuc. The system has an orbital period of $29.5~d$ and the orbit is slightly eccentric with $e=0.056$. We obtain $M_p=0.8762\\pm 0.0048\\,M_\\odot$, $R_p=1.3148\\pm 0.0051\\,R_\\odot$, $L_p=1.94\\pm 0.21\\,L_\\odot$ for the primary  and $M_s=0.8588\\pm 0.0060\\,M_\\odot$, $R_s=1.1616\\pm 0.0062\\,R_\\odot$, $L_s=1.53\\pm 0.17\\,L_\\odot$ for the secondary. These components of V69 are the first Population~II stars with masses and radii derived directly and with an accuracy of better than 1\\%. We measure an apparent distance modulus of $(m-M)_{V}=13.35\\pm 0.08$ to V69. We compare the absolute parameters of V69 with five sets  of stellar evolution  models and estimate the age of V69 using mass-luminosity-age, mass-radius-age, and turnoff mass - age relations. The masses, radii, and luminosities of the component stars are determined well enough that the measurement of ages is dominated by systematic differences between the evolutionary models, in particular, the adopted helium abundance. By comparing the observations to Dartmouth model isochrones we estimate the age of V69 to be 11.25$\\pm0.21$(random)$\\pm$0.85(systematic)~Gyr assuming [Fe/H]~=~-0.70, [$\\alpha$/Fe]~=~0.4, and  $Y$~=~0.255. The determination of the distance to V69, and hence to 47~Tuc, can be further improved when infrared eclipse photometry is obtained for the variable. ",
        "introduction": "\\label{intro}  Detached eclipsing double-line binary stars are the fundamental astrophysical laboratory for the determination of stellar parameters of mass and radius. Luminosities can be derived using measured parallaxes or from empirical color -- effective temperature relations. These data are the fundamental tests of stellar evolution models. Many field Population~I systems are known at solar mass and larger \\citep{andersen} and modern high accuracy measurements of masses, luminosities, and radii of the component stars are in general agreement with evolution models \\citep[see, for example,][]{lacy05,lacy08, clausen08}. Similar results are obtained for studies of individual binaries in the old open clusters NGC~188 \\citep{meibom09}, NGC~2243 \\citep{kaluzny06}, and NGC~6791 \\citep{grundahl08}. Recent wide field photometric surveys have identified numerous low mass systems, and for these K and M stars the common theme of a comparison of component properties with evolution models is that the models systematically underestimate the radii of the components in these binaries.  Summaries of recent measurements can be found in \\citet{lopez06}, \\citet{lopez07}, and \\citet{blake08}.  The situation is even less clear for Population~II stars. With the exception of CM Dra, for which the component masses are $\\sim$0.2$\\,M_\\odot$ \\citep{lacy77}, and the $\\omega$~Cen binary OGLEGC-V17, for which the analysis is compromised by an uncertain determination of the metallicity \\citep{thompson01,kal02}, there are no known Population~II detached double-line eclipsing binaries with main-sequence components.  \\citet{torres02} used interferometric observations of HD~195987 ([Fe/H] $\\sim$ -0.6) to derive an orbit and measure the masses of the components. The radiative properties of the two components agree with a suite of models given some slight modifications of input parameters. While direct measurements of the radii of the components are not possible because HD~195987 is not an eclipsing binary, estimates of the radii can be derived from the orbital parallax, the bolometric flux, and the estimated effective temperatures. Here again the measured radii are larger than the models by some ten per cent. \\citet{boyajian08} have measured the angular diameter of the G subdwarf $\\mu$~Cas ([Fe/H $\\sim$ -0.8). This measurement provides a radius when combined with the Hipparchos parallax for $\\mu$~Cas. For this star the models underpredict this measured radius by about five per cent. The masses of the components of $\\mu$~Cas are only known to about ten per cent, so the model comparisons here are not well constrained.  There is a clear need to locate and study Population~II detached eclipsing binary stars to obtain accurate masses and radii of their component stars. Stellar evolution models are becoming increasingly sophisticated and are used to fit  observed cluster color-magnitude diagrams (CMD). The cluster CMD's are themselves improving in quality, with homogeneous surveys being conducted with the Hubble Space Telescope \\citep[see, for example, ][]{sarajdeni07}. Careful empirical tests of these models are an essential next step.    This is the first paper in a series devoted to the study of detached eclipsing double-line binaries (DEB) in Galactic globular clusters with components on the cluster main sequence or subgiant branch. The Cluster AgeS Experiment (CASE) has the goal of determining the basic stellar parameters (masses, luminosities, and radii) of the components of cluster binaries to a precision of better than  1\\% in order to measure cluster ages and distances, and to test stellar evolution models. The methods and assumptions utilize basic and simple  approaches offered by the field of eclipsing double-line spectroscopic binaries as described in \\citet{paczynski97} and \\citet{thompson01}. Previous CASE papers have discussed blue straggler systems in $\\omega$~Cen \\citep{kal07a} and 47~Tuc \\citep{kal07b}, and an SB1 binary in NGC~6397 \\cite{kal08}.   The eclipsing binary  V69-47~Tuc (hereinafter V69) was discovered by \\citet{weldrake} during a survey for variable stars in the field of the globular cluster 47~Tuc. They presented an $I$-band light curve for the variable and proposed an orbital period of $P=5.229$~d. The light curve phased with this period shows only one eclipse. The variable is located at the top of main sequence in the cluster color-magnitude diagram, and thus is of potential great interest for measurements of the cluster age and distance.  In this paper we report the  results of photometric and spectroscopic observations aimed at a determination of the absolute parameters of the components of V69. Section~\\ref{phot} describes the photometry of the variable and the determination of an orbital ephemeris. Section~\\ref{spec} presents the radial velocity observations. The combined photometric and spectroscopic element solutions are given in Section~\\ref{comb} while the membership in 47~Tuc is discussed in Section~\\ref{memb}. In Section~\\ref{age} we compare the properties of the components of V69 to a selection of stellar evolution models with an emphasis on estimating the age of the system. Finally in  Section~\\ref{disc} we summarize our findings. ",
        "conclusions": "\\label{disc} To the best of our knowledge, these measurements of the masses and radii of the  components of V69 are the first such high accuracy (better than 1\\%) measurements to be made for Population II stars. The binary is a member of the globular cluster 47~Tuc and so the determination of its distance and age applies to the cluster as well. We obtained a distance modulus of $(m-M)_{V}=13.35\\pm 0.08$. The main source of error in the distance estimate is the calibration dependent estimate of effective temperatures which we used to derive the bolometric luminosities for the component stars.   A comparison of the measured masses, luminosities, and radii of the components to stellar evolution models suggests that the age of the system and hence the globular cluster 47~Tuc can be measured to a statistical accuracy of about 0.25 Gyr. However, it is important to understand the assumptions that go into any one model. In particular, the derived ages are very sensitive to the adopted helium abundance. We derive an age for 47~Tuc of 11.25$\\pm0.21$$\\pm$0.45~Gyr for a helium abundance $Y$~=~0.255 using Dartmouth model isochrones. All models give similar ages when the effects of helium abundance are taken into account.  Comparison of Dartmouth evolutionary tracks calculated for the measured masses of the primary and secondary indicate that the helium abundance can be measured to an accuracy of about 0.03 for each of the components. We estimate a helium abundance of $Y = 0.269^{+0.017}_{-0.021}$ for 47~Tuc. The measured masses, radii, and luminosities of the components of V69 are consistent with Dartmouth models assuming [Fe/H] = -0.71, [$\\alpha$/Fe] = + 0.4, and Y = 0.27.   The radii of both stars are known with high accuracy, and it is therefore  possible to obtain a more accurate and robust distance determination based on the surface brightness method \\citep{barnes76,lacy77,thompson01}. The empirical calibration of surface brightness relations for dwarf and subgiant stars is improving \\citep{dibenedetto98,kervella,buermann06}, and it is reasonable to imagine that a distance accurate to a few per cent can be measured with accurate radii and $(V-K)$ colors. We are in the process of collecting near IR eclipse profile photometry of both V69 and V228 \\citep{kal07b} in order to measure the distances to these two binary stars in this way. These data will improve the estimates of the bolometric luminosities of the components and lead to a more accurate measurement of the helium abundance and hence the absolute age of 47 Tuc.  Finally, we note that contributions to the errors in the radii are dominated by the photometric solution. Given the large inclination, the errors in the masses are completely dominated by the orbital solution. An identical doubling of the existing set of radial velocity observations leads to a 33$\\%$ improvement in the mass estimates, and a subsequent improvement in the age estimates. The system is relatively bright, and the prospects are good that a substantial improvement in the measured masses can be achieved with further radial velocity observations."
    },
    "0910/0910.1453_arXiv.txt": {
        "abstract": "We present an analysis of timing residual (noise) of $54$ pulsars obtained from 25-m radio telescope at Urumqi Observatory with a time span of $5\\sim 8$ years, dealing with statistics of the Hurst parameter. The majority of these pulsars were selected to have timing noise that look like white noise rather than smooth curves. The results are compared with artificial series of different constant pairwise covariances. Despite the noise like appearance, many timing residual series showed Hurst parameters significantly deviated from that of independent series. We concluded that Hurst parameter may be capable of detecting dependence in timing residual and of distinguishing  chaotic behavior from random processes. ",
        "introduction": "All pulsars show a remarkable uniformity of rotation rate on a time scale of a few days as expected of an isolated spinning body with large stable moment of inertia. \\citep{Lyne06} The angular momentum of radio pulsar is slowly decreasing through slowdown torque of the magnetic dipole radiation. However, some very interesting irregularities in pulsar rotation have been observed which are termed as {\\em timing noise}.  It is anticipated that valuable information of many interesting physical processes related to pulsars is coded in the timing noises, thus employing statistical measures to characterize timing noises is important to the study of pulsars and thus the properties of the matter at supra-nuclear densities. Efforts to quantify timing noise have been tried as early as timing noise was firstly recognized, for instance according to random walk of different quantities \\citep{Boynton72}, most current models such as vortex creeping are still restricted to treatment of timing noise only as random process in certain quantities. One exception was presented by \\citep{Harding90}, who analyzed timing data of Vela pulsar to look for evidence of chaotic behavior by ``correlation sum'' technique to estimate fractal dimension of the system. However, despite possible suggestions they concluded that ``correlation sum'' estimator may be unable to distinguish between random and chaotic processes.  However, any statistical representation of data has their own biases, employing a large types of statistical measures is essentially vital to a fair understanding of the timing noises. Furthermore, the number of observed pulsars has accumulated to $\\sim 10^3$ \\citep{Manchester06} but collected data is often incomplete for a conclusive analysis, it is therefore crucial to diagnose current available but limited data, the results of which could guide us in future observation to concentrate on those pulsars with anomalous timing noise. Here we are introducing a statistical method rarely used in time domain astronomy -- the Hurst parameter analysis, which is actually sensitive to the type of the inherent correlation among the time series. Our practical analysis of the timing data observed by the 25-meter radio telescope at Urumqi Observatory of $54$ pulsars indicates that Hurst parameter analysis might be capable of detecting anomalous signals which disguise themselves as noises. ",
        "conclusions": "\\label{sec_cncl} We calculated Hurst parameter of $54$ radio pulsars with white-noise-like timing residual obtained from Nanshan telescope and compared the results with artificial (anti-)persistent series. The majority of pulsars from our selection have Hurst parameters around $0.5$ and not far from Hurst parameter calculated for independent series. However, we found 9 pulsars (PSRs J0147$+$5922, J0357$+$5236, J0612$+$3721, J0630$-$2834, J0823$+$0159 and J0837$-$4135 showing persistent trend and PSRs J0055$+$5117, J1022$+$1001, and J1842$-$0359 showing anti-persistent trend) with interesting $H$ values despite having white-noise-like timing residual. Comparison with artificial series confirm that these trends cannot be attributed entirely to effects caused by ROS algorithm or large uncertainty in timing residual. This shows that our algorithm may be capable of detecting hidden correlation in apparently noise-like timing residual. We therefore suggest that these 9 pulsars be monitored continuously to confirm or disprove long-range dependence and search for possible physical process behind such correlation.  As for the selection, we picked $54$ pulsars with relatively small period derivative $\\dot{P}$, and white-noise-like timing residual rather than the typical red-noise pattern of the three kind of noise model (PN, FN and SN) first proposed by Boynton et al. (1972) and of course we made no attempt at doing correlation between Hurst parameter and other pulsar properties nor obtaining statistics of Hurst parameter values for large population. Our method is aimed at finding possibility that timing noise that resembles white-noise is not really generated by a random process. As is well known there are pulsars with smooth timing noise that differ largely from white noise (such as J0332+5434, J0406+6138, J0826+2637, etc.) and they would yields quite large Hurst parameter if calculated using our algorithm and some even exceed $1$. In that case we do not need the deviation from $0.5$ to conclude the obvious dependent nature of timing noise and other method should be used to analyze possible chaotic nature of these timing noise series.  There are various physical processes that might be responsible for the long-range dependence of pulsar timing noise. These can be classified into three groups: from interior of neutron star, i.e. due to fluctuation of internal (e.g. micro-quake due to partial release of elastic energy \\cite{Pines72} and random pinning and unpinning of vortex lines \\cite{Packard72}, \\cite{Anderson75}) and external (e.g. accretion flow \\cite{Lamb78}) torques; from emission process (e.g. magnetosphere activities \\cite{Cheng87}) and from propagation of radio emission. For the last class of origin, it has long been proposed that timing noise can be used to set upper limits on gravitational wave (\\cite{Bertotti83}) and that given enough time pulsar timing array might be the first equipment to directly detect gravitational wave (\\cite{Manchester062}). All of these (including gravitational wave detection which aims at detecting random background) predict randomness in evolution of certain quantity while Hurst parameter is capable of uncovering chaotic behavior that is hidden in timing noise. The physical origin that leads to dependence series is much more limited than random fluctuations. Therefore, long-term dependence detected by Hurst parameter might reveal more detailed information about the physical origin of timing noise. For instance it has been proposed that under certain conditions Euler equations for rotating object with magnetic dipole moment misaligned with rotation axis would show chaotic spin-down behavior (see \\cite{Harding90}). By definition any dynamical process can exhibit chaotic behavior and thus dependence in long-term timing noise data if the system is governed by differential equations whose solution is highly sensitive to initial conditions. We expect that the application of Hurst parameter to timing noise of longer time span reveal more evidence for new physics."
    },
    "0910/0910.4937_arXiv.txt": {
        "abstract": "{Sub-millimeter and Far-IR observations have shown the presence of a significant amount of warm (few hundred K) and dense ($n(\\rm H_2)\\ge10^4~\\3cm$) gas in sources ranging from active star forming regions to the vicinity of the Galactic center. Since the main cooling lines of the gas phase are important tracers of the interstellar medium in Galactic and extragalactic sources, proper and detailed understanding of their emission, and the ambient conditions of the emitting gas, is necessary for a robust interpretation of the observations. } {With high resolution ($7''-9''$) maps ($\\sim3\\times3$ pc$^2$) of mid-$J$ molecular lines we aim to probe the physical conditions and spatial distribution of the warm (50 to few hundred K) and dense gas ($n(\\rm H_2)>10^5~\\3cm$) across the interface region of the nearly edge-on M17 SW nebula. } {We have used the dual color multiple pixel receiver CHAMP$^+$ on APEX telescope to obtain a $5'.3\\times4'.7$ map of the $J=6\\rightarrow5$ and $J=7\\rightarrow6$ transitions of \\twco, the \\thco $J=6\\rightarrow5$ line, and the $^3P_2\\rightarrow{^3P_1}$ 370 $\\mu$m fine-structure transition of \\ci in M17 SW. LTE and non-LTE radiative transfer models are used to constrain the ambient conditions. } {The warm gas extends up to a distance of $\\sim2.2$ pc from the M17 SW ridge. The \\thco\\ $J=6\\rightarrow5$ and \\ci\\ 370 $\\mu$m lines have a narrower spatial extent of about 1.3 pc along a strip line at P.A=$63^{\\circ}$. The structure and distribution of the \\ci $^3P_2\\rightarrow{^3P_1}$ 370 $\\mu$m map indicate that its emission arises from the interclump medium with densities of the order of $10^3~\\3cm$. } {The warmest gas is located along the ridge of the cloud, close to the ionization front. An LTE approximation indicates that the excitation temperature of the embedded clumps goes up to $\\sim120$ K. The non-LTE model suggests that the kinetic temperature at four selected positions cannot exceed 230 K in clumps of density $n(\\rm H_2)\\sim5\\times10^5~\\3cm$, and that the warm ($T_k>100$ K) and dense ($n(\\rm H_2)\\ge10^4~\\3cm$) gas traced by the mid-$J$ \\twco lines represent just about 2\\% of the bulk of the molecular gas. The clump volume filling factor ranges between 0.04 and 0.11 at these positions. } ",
        "introduction": "The heating and cooling balance in photon-dominated regions (PDRs) remains an active study of research. The comprehensive understanding of PDRs requires observations of large areas close to radiation sources, and of a wide wavelength range covering various emissions of atoms, molecules, and grains. In particular, mid-$J$ CO lines have been detected in almost all known massive Galactic star forming regions (e.g. Orion Nebula, W51, Cepheus A, NGC 2024). This indicates that warm ($T_K\\ge50$ K) and dense ($n(H_2)\\ge10^4~\\3cm$) gas is common, and probably of importance in most OB star forming regions. The mid-$J$ CO lines detected in regions like, e.g. M17, Cepheus A and W51, have relatively narrow line widths of 5--10 \\kms, although not as narrow as the line widths observed in cold quiescent cloud cores.  Observations of the $J=6\\rightarrow5$ and $J=7\\rightarrow6$ transitions of \\twco in several massive star forming regions indicate that the warm emitting gas is confined to narrow ($<1$ pc) zones close to the ionization front. These observations favor photoelectric heating of the warm gas by UV radiation fields outside the HII regions (e.g. Harris \\etal\\ 1987; Graf \\etal\\ 1993; Yamamoto \\etal\\ 2001; Kramer \\etal\\ 2004 and 2008). Nevertheless, shocks may also be an important source of heating in high velocity wing sources like, Orion, W51 and W49 (Jaffe \\etal\\ 1987).  Because of its nearly edge-on geometry, and the large amount of observational data available in the literature, M17 SW is one of the best Galactic regions to study the entire structure of PDRs from the exciting sources to the ionization front, and the succession (or not) of H$_2$, \\ci and CO emissions, as predicted by PDR models (Icke, Gatley, \\& Israel 1980; Felli, Churchwell, \\& Massi 1984; Meixner et al. 1992; Meijerink \\& Spaans 2005). M17 SW is also one of the few star forming regions for which the magnetic field strength can be measured in the PDR interface, and where the structure of the neutral and molecular gas seems to be dominated by magnetic pressure rather than by gas pressure (Pellegrini \\etal\\ 2007).  M17 SW is a giant molecular cloud at a distance of 2.2 kpc, illuminated by a highly obscured ($A_v>10$ mag) cluster of several OB stars (among $\\ga100$ stars) at about 1 pc to the East (Beetz \\etal\\ 1976; Hanson \\etal\\ 1997). It also harbours a number of candidate young stellar objects that have recently been found (Povich \\etal\\ 2009). Several studies of molecular emission, excitation and line profiles (e.g. Snell \\etal~1984; Martin, Sanders \\& Hills 1984; Stutzki \\& G\\\"usten 1990) from the M17 SW core indicate that the structure of the gas is highly clumped rather than homogeneous. Emission of \\ci and \\cii was detected more than a parsec into the molecular cloud along cuts through the interface region (Keene \\etal\\ 1985; Genzel \\etal\\ 1988; Stutzki \\etal\\ 1988). These results, as well as those found in other star forming regions like S106, the Orion Molecular Cloud, and the NGC~7023 Nebula (e.g. Gerin \\& Phillips 1998; Yamamoto \\etal\\ 2001; Schneider \\etal\\ 2002, 2003; Mookerjea \\etal\\ 2003) do not agree with the atomic and molecular stratification predicted by standard steady-state PDR models. However, the extended \\ci $^3P_1\\rightarrow{^3P_0}$ and \\thco $J=2\\rightarrow1$ emission in S140 have been successfully explained by a stationary, but clumpy, PDR model (Spaans 1996; Spaans \\& van Dishoeck 1997). Hence, the lack of stratification in \\ci, \\cii and CO is a result that can be expected for inhomogeneous clouds, where each clump acts as an individual PDR. On the other hand, a partial face-on illumination of the molecular clouds would also suppress stratification.  Based on analysis of low-$J$ lines of \\twco, \\thco and CH$_3$CCH data, the temperature towards the M17 SW cloud core has been estimated as 50--60 K, whereas the mean cloud temperature has been found to be about 30--35 K (e.g. G\\\"usten \\& Fiebig 1988; Bergin \\etal\\ 1994; Wilson \\etal\\ 1999; Howe \\etal\\ 2000; Snell \\etal\\ 2000). Temperatures of $\\sim275$ K has been estimated from NH$_3$ observations (G\\\"usten \\& Fiebig 1988) towards the VLA continuum arc, which agree with estimates from highly excited \\twco transitions (Harris \\etal\\ 1987). Multitransition CS and HC$_3$N observations indicates that the density at the core region of M17 SW is about $6\\times10^5~\\3cm$ (e.g., Snell \\etal\\ 1984; Wang \\etal\\ 1993; Bergin, Snell \\& Goldsmith 1996). While densities up to $3\\times10^6~\\3cm$ have been estimated towards the north rim with multitransition observations of NH$_3$, which indicates that ammonia is coexistent with high density material traced in CS and HCN (G\\\"usten \\& Fiebig 1988). The UV radiation field $G_0$ has been estimated to be of the order of $10^4$ in units of the ambient interstellar radiation field ($1.2\\times10^{-4}~\\rm ergs~s^{-1}~cm^{-1}~sr^{-1}$, Habing 1968; Meixner \\etal\\ 1992).   However, most of the millimeter-wave molecular observations in M17 SW are sensitive only to low temperatures ($<100$ K), and the few available data of mid-$J$ CO and \\ci lines (consisting mostly of cuts across the ionization front and observations at few selected positions) are limited in spatial resolution and extent (e.g. Harris \\etal\\ 1987; Stutzki \\etal\\ 1988; Genzel \\etal\\ 1988; Stutzki \\& G\\\"usten 1990; Meixner \\etal\\ 1992; Graf \\etal\\ 1993; Howe \\etal\\ 2000). Therefore, in this work we present maps ($\\sim3\\times3$ pc$^2$) of mid-$J$ molecular (\\twco and \\thco) and atomic (\\ci) gas, with excellent high resolution ($9.4''-7.7''$), which advances existing work in M17 SW.  The observations were done with CHAMP$^+$ (Carbon Heterodyne Array of the MPIfR) on the Atacama Pathfinder EXperiment (APEX\\footnote{This publication is based on data acquired with the Atacama Pathfinder Experiment (APEX). APEX is a collaboration between the Max-Planck-Institut fur Radioastronomie, the European Southern Observatory, and the Onsala Space Observatory}) (G\\\"usten \\etal\\ 2006). The multiple pixels at two submm frequencies of CHAMP+, allow for efficient mapping of $\\sim$arcmin regions, and provide the ability to observe simultaneously the emission from the $J=6\\rightarrow5$ and $J=7\\rightarrow6$ rotational transitions of \\twco at 691.473 GHz and 806.652 GHz, respectively. We also observed the $J=6\\rightarrow5$ transition of \\thco at 661.067 GHz and the $^3P_2\\rightarrow~^3P_1$ 370 $\\mu$m (hereafter: $2\\rightarrow1$) fine-structure transition of \\ci at 809.342 GHz.  Since the gas phase cools mainly via the atomic fine structure lines of \\oi, \\cii, \\ci, and the rotational CO lines (e.g. Kaufman \\etal\\ 1999, Meijerink \\& Spaans 2005), these carbon bearing species presented here are very important coolants in the interstellar medium (ISM) of a variety of sources in the Universe, from Galactic star forming regions, the Milky Way as a galaxy, and external galaxies up to high redshifts (e.g. Fixsen \\etal\\ 1999; Weiss \\etal\\ 2003; Kramer \\etal\\ 2005; Bayet \\etal\\ 2006; Jakob \\etal\\ 2007).  The case of M17 SW can be considered as a proxy for extra galactic star forming regions. M17 SW is not special, nor does it need to be, compared to other massive star-forming regions like Orion, W49, Cepheus A, or W51. Still, it does allow feedback effects, expected to be important for starburst and active galaxies, to be studied in great spatial detail. Comparison of the local line ratios to extra-galactic regions can then shed light on the properties of massive star forming regions that drive the energetics of active galaxies. Our results will be of great use for future high resolution observations, since molecular clouds of the size of the maps we present will be resolved by ALMA at the distance ($\\sim14$ Mpc) of galaxies like NGC~1068.  The main purpose of this work is to explore the actual spatial distribution of the mid-$J$ \\twco\\ and \\ci lines in M17 SW, and to test the ambient conditions of the warm gas. A simple LTE model based on the ratio between the \\twco and \\thco $J=6\\rightarrow5$ lines is used to probe the temperature of the warm ($T_K\\sim100$ K) and dense ($n_{\\rm H}>10^5~\\3cm$) molecular gas. Then a non-LTE model is used to test the ambient conditions at four selected positions. In a follow up work we will present an elaborate model of these high resolution data.  The most frequent references to Stutzki \\etal\\ (1988), Stutzki \\& G\\\"usten (1990) and Meixner \\etal\\ (1992) will be referred to as S88, SG90 and M92, respectively. The organization of this article is as follows. In \\S2 we describe the observations. The maps of the four lines observed are presented in \\S3. The modelling and analysis of the ambient conditions are presented in \\S4. And the conclusions and final remarks are presented in \\S5. ",
        "conclusions": "We have used the dual color heterodyne receiver array of 7 pixels CHAMP$^+$ on the APEX telescope, to map a region of about 2.6 pc $\\times$ 2.9 pc in the $J = 6\\rightarrow5$ and $J = 7\\rightarrow6$ lines of \\twco, the \\thco $J = 6\\rightarrow5$ and the $^3P_2\\rightarrow{^3P_1}$ 370 $\\mu$m ($J=2\\rightarrow1$) fine-structure transition of \\ci in M17 SW nebula.  The completely different structure and distribution of the $^3P_2\\rightarrow{^3P_1}$ 370 $\\mu$m emission, and its critical density, indicate that this emission arises from the interclump medium ($\\sim3\\times10^3~\\3cm$). On the other hand, the mid-$J$ lines of \\twco and the isotope emissions, arise from the high density ($\\sim5\\times10^5~\\3cm$) and clumpy region.  The spatial extent of the warm gas (40-230 K) traced by the \\twco $J=7\\rightarrow6$ line is about 2.2 pc from the ridge of the M17 SW complex, which is smaller than the extent observed in the low-$J$ \\twco and C$^{18}$O lines reported in previous work. The \\thco\\ $J=6\\rightarrow5$ and \\ci\\ 370 $\\mu$m lines, have a narrower spatial extent of about 1.3 pc along a strip line at P.A=$63^{\\circ}$.  An LTE approximation of the excitation temperature provides lower limits for the kinetic temperature. The warmest gas is located along the ridge of the cloud, close to the ionization front. In this region the excitation temperature range between 40 and 120 K. A non-LTE estimate of the ambient conditions at four selected positions of M17 SW indicates that the high density clumps ($\\sim5\\times10^5~\\3cm$) cannot have temperatures higher than 230 K. The warm ($T_k>100$ K) and dense ($n(\\rm H_2)\\ge10^4~\\3cm$) gas traced at the four selected positions by the mid-$J$ \\twco lines represent $\\sim2$\\% of the bulk of the molecular gas traced by the low-$J$ \\twco lines. Volume filling factors of the warm gas ranging from 0.04 to 0.11 were found at these positions."
    },
    "0910/0910.4330_arXiv.txt": {
        "abstract": "% There is no shortage of energy around to solve the overcooling problem of cooling flow clusters. AGNs, as well as gravitational energy are both energetic enough to balance the cooling of cores of clusters. The challenge is to couple this energy to the baryons efficiently enough, and to distribute the energy in a manner that will not contradict observational constraints of metalicity and entropy profiles. Here we propose that if a small fraction of the baryons that are accreted to the cluster halo are in the form of cold clumps, they would interact with the hot gas component via hydrodynamic drag. We show that such clumps carry enough energy, penetrate to the center, and heat the core significantly. We then study the dynamic response of the cluster to this kind of heating using a 1D hydrodynamic simulation with sub-grid clump heating, and produce reasonable entropy profile in a dynamic self-consistent way. ",
        "introduction": "Galaxy clusters grow by accreting dark matter and baryons from their surroundings. This is partly a continuous process of relatively smooth accretion, and partly via mergers. The existence of a visible, extended X-ray halo as well as limits on the size of the CD galaxy are two indicators that the smooth, continuous accretion is the more pronounced of the two. The baryons that are accreted in this process settle in hydrostatic equilibrium within the cluster's potential well. During this process they achieve virial equilibrium  by passing through the virial shock and converting the kinetic energy of the infall to thermal energy. The baryons, at a temperature of a few $keVs$, cool primary by emitting Bremsstrahlung radiation which easily escape the halo and is observed by X-ray telescopes. After de-projection (tomography) of the luminosity and spectrum, radial profiles of temperature and densities are derived, and one can deduce the cooling times of the gas, which, for the centers of cooling flow cluster, is $\\le 1Gyr$, much shorter than the age of typical clusters that, according to $\\Lambda CDM$ formation history \\citep{press74} should have been in place at $z \\ge 1$. Had this gas cooled, \\emph{i)} cool gas would have been seen in the halo (no gas below $T=\\frac{1}{3}T_{vir}$ is observed), \\emph{ii)} a census of all the baryons in galaxies is significantly smaller than amount of gas that was expected to cool from the halo, and \\emph{iii)} the star formation within the CD would have been $10^2-10^3M_\\odot/yr$ (two orders of magnitude larger than typically observed). These three contradictions are three manifestations of the overcooling problem of cooling flow clusters. It is highly unlikely that there is a ``hidden'' baryonic component in cluster halos, so most explanations invoke some kind of heating mechanisms that would balance the cooling, keeping the gas hot and diffuse.  The energies needed to compensate for the rapid cooling of cluster halos are $\\sim 10^{45}erg/sec$ which, over the lifetime of the cluster amounts to $\\sim 10^{62} erg$. These required energy rates can originate from AGN emission  from the CD \\citep[][as well as many others]{ciotti97}, and by gravitational energy. Diffuse baryons falling into a gravitational well convert gravitational energy to kinetic, and ultimately to thermal. This thermal input into the system is usually local, and acts to heat the infalling gas itself. It is necessary to couple this freshly accreted, hot gas, with the central cooling gas. \\citet{narayan01} have studied conduction and turbulence that could potentially couple the external part to the inner halo. \\citet{kim03}, as well as others, deduce that although the amount of energy is sufficient, the conduction coefficient is not enough and turbulence will produce the wrong entropy and metalicity profiles.  In this proceedings, we shall present a novel mechanism of gravitational heating of the central halo gas by the hydrodynamic interactions between cold clumps of accreted gas and the halo gas \\citep{db08} and \\citetext{Birnboim \\& Dekel 2009, in preparation}. The structure of this proceeding paper is as follows: First (section \\ref{sec:energetics}) we show that the accreting gas carries a sufficient rate of energy to compensate for the cooling. Then we model infalling cold clumps, and study the physical processes of their interaction with the diffuse gas, instabilities and survivability (section \\ref{sec:physics}). Using this model, we study the valid parameter space of these clumps (section \\ref{sec:monte}). We further use these insights to construct a sub-grid model for 1D hydrodynamic simulation and present simulations of a cluster with and without such clumps (section \\ref{sec:hydro}). Finally we discuss possible origins of these clumps, and summarize (section \\ref{sec:summary}). ",
        "conclusions": "\\label{sec:summary} We have shown that a sufficient amount of energy is released in the gravitational accretion to solve the overcooling problem of clusters (larger than $M_{vir}\\ge 10^{13}M_\\odot$), if we assume that some part of the baryons penetrates to the inner parts rather that being stopped at the virial radius. The coupling between the incoming accreted mass and the ambient gas is achieved by assuming that the accreted baryons includes gaseous cold clump, that penetrate through the virial shock and heat the gas via hydrodynamic drag. A parameter space survey indicates that for halos larger than $10^{13}M_\\odot$ and clump masses in the range $10^5\\le m_c \\le 10^8M_\\odot$, clump heating can potentially solve the overcooling problem. The dynamic response of cluster halos to clump heating, cooling and mass deposition in the inner parts in then examined. For the set of values $M_{vir}=3 \\cdot 10^{14}M_\\odot$, $m_c=10^7M_\\odot$, $f_b+f_c=0.1$, $f_c=0.01$, the resulting entropy and temperature profiles match typical observed cooling flow clusters reasonably well. The physical processes discussed here are already included in 3D hydrodynamic simulations but we note that to simulate the drag and the formation of such clumps each $10^7M_\\odot$ clump need to be well resolved, which is typically not the case for $10^{15}M_\\odot$ clusters."
    },
    "0910/0910.5336_arXiv.txt": {
        "abstract": "The X-3.4 class flare of 13 December 2006 was observed with a high cadence of 2 minutes at 0.2 arc-sec resolution by HINODE/SOT FG instrument. The flare  ribbons could be seen in G-band images also. A careful analysis of these observations after proper registration of images show  flare related changes in penumbral filaments of the associated sunspot, for the first time. The observations of sunspot deformation, decay of penumbral area and changes in magnetic flux during large flares have been reported earlier in the literature. In this {\\it Letter}, we report lateral motion of the penumbral filaments in a sheared region of the $\\delta$-sunspot during the X-class flare. Such shifts have not been seen earlier. The lateral motion occurs in two phases, (i) motion before the flare ribbons move across the penumbral filaments and (ii) motion afterwards. The former motion is directed away from expanding flare ribbons and lasts for about four minutes.  The latter motion is directed in the opposite direction and lasts for more than forty minutes. Further, we locate a patch in adjacent opposite polarity spot moving in opposite direction to the penumbral filaments. Together these patches represent conjugate foot-points on either side of the polarity inversion line (PIL), moving towards each other. This converging motion could be interpreted as shrinkage of field lines. ",
        "introduction": "Photospheric changes during large flares have been studied for a long time \\citep{Giovanelli40,Howard63,Rust76,Tanaka78}. Such changes have been observed successfully in various parameters like, shear angle \\citep{Ambastha93,Wang94,Hagyard99}, sunspot morphology \\citep{Anwar93}, magnetic flux \\citep{Lara00, Spirock02, Wang02a, Wang02b, Sudol05, Zharkova2005}, penumbral size \\citep{Wang04,Deng05, Liu05}. Most of these observed changes are located in the penumbra of the sunspot. \\cite{Hagyard99} argued that penumbral regions contribute most to magnetic energy of a force-free field (Low 1985) given by $E_M= \\int{(x B_x + y B_y)B_z dx dy}$. Since $B_z=0$ near Polarity Inversion Line (PIL) and $B_T = \\sqrt{B_x^2 + B_y^2} \\approx 0$ in the umbra, these regions contribute little to magnetic energy of the force-free field. Since, flares are driven by magnetic energy changes, the most promising locations to look for flare related changes are the regions between umbra and PIL. This could explain why most of the flare related studies show changes in the penumbra.  \\cite{Hagyard99} found interesting coherent variation of the shear angle in patches of the magnetic field on timescales shorter than the photospheric Alfven travel time over the patches. They had conjectured that electrodynamic effects of coronal magnetic field variations could induce coherent magnetic variations even at the photospheric levels. They also concluded that instruments with greater sensitivity could throw more light on these induced field variations. The Solar Optical Telescope (SOT; \\cite{Tsuneta08}) onboard Japanese space mission {\\it Hinode} \\citep{Kosugi07} is one such instrument capable of detecting these changes.  The X-3.4 class flare, observed in NOAA 10930 on 13 December 2006 during 02:00 UT to 03:00 UT, by {\\it Hinode} spacecraft, provided such an opportunity.  Various aspects of this active region have been reported. Three dimensional magnetic field reconstruction of this region was performed by \\cite{Schrijver08} using  NLFFF extrapolation methods. They found regions of strong vertical current density between the two spots which  decreased significantly after the flare. Further, the free energy also decreased after the flare. The variation  of  the  weighted mean shear and total magnetic shear \\citep{Wang08}  with altitude was studied using extrapolated fields by \\cite{Jing08}. They found that near the PIL, up to a height of  $\\approx 8 Mm$ the shear increases after the flare and decreases beyond that  up to $\\approx 70 Mm$. Beyond 70 Mm the field remains close to potential before and after the flare.  However, the cadence of these vector magnetograms is insufficient to search for similar flare related changes as observed by \\cite{Hagyard99}.   In order to detect flare related changes in  magnetized regions like sunspots one can also use high angular resolution images obtained at a high-cadence,  during the flare interval. Such observations are nowadays routinely available from Filtergraph (FG; \\cite{Tsuneta08}) instrument onboard {\\it Hinode} space mission. The high resolution images allow us to investigate the evolution of the fine structure like penumbral filaments during a flare, e.g lateral motion and/or twist. In this {\\it Letter}, we use  FG observations from {\\it Hinode} to investigate  flare-induced changes during a X-class flare.   In section 2.1 we describe the observations while the method of registration applied on the data to establish the coherent motions are described in section 2.2. We present our results in section 3 while the discussion and conclusions are given in section 4. ",
        "conclusions": "In this {\\it Letter}, we present sub-arcsec (0.1\"/pixel) observations in G-band at a cadence of 2 minutes, taken during 13 December 2006 X3.4 flare. High-resolution observations of the photosphere during a large X-class flare have never been done before with the kind of stability and image quality that HINODE/SOT provides. A careful analysis,  after proper registration of these images, show  flare related changes in penumbral filaments of the associated sunspot, for the first time, in the form of coherent lateral motions of the penumbral filaments.  We found the motion of the penumbral filaments to exist in two phases: (i) the first phase is short-lived (about 4 minutes) with the time of maximum displacement coinciding with the peak of the microwave flux, and (ii)  the second phase is of longer duration (more than 40 minutes).  Further, we observe the patches corresponding to conjugate foot-points in either spots moving towards each other. These conjugate foot-points, labeled as `1' and `2' and marked by arrows, are clearly seen in the online G-band movie. This converging motion of conjugate foot-points could be interpreted as shrinkage of field lines during flare, which is conjectured by \\cite{Hudson00}. However, we lack high-cadence vector field measurements in this case and our present inferences are purely from G-band images. We have also shown, in section 3.4, a decrease in the global twist after the flare. It will be interesting to see the evolution of local twist as well as global twist during the flare. We expect to detect flare-related vector field changes with  high-cadence vector magnetograms from the upcoming SDO mission.  \\cite{Schrijver08} found strong arching filamentary current system embedded in the magnetic flux emerging between the penumbrae of the two spots, prior to flare.  These currents are significantly reduced after the flare. This current system seems co-spatial with the field lines that connect the conjugate foot-points discussed above.  Currents in astrophysical plasmas are produced by the distortions of the magnetic fields given by the curl of the field \\citep{Parker79}. Generally these distortions are driven by plasma dynamics. The coherent nature of the distortions in the present observations, occurring within a time scale shorter than the Alfven travel time over the affected area, rules out the plasma dynamics of the photosphere as the cause of the distortions. It is more likely that the distortions are caused by the fields induced in the photosphere  in response to the restructuring  (like untwisting of the global field, as inferred from the decrease in SSA) of the coronal magnetic fields,  which are line tied to the photosphere. We have estimated the energy of the induced lateral motions of the penumbral filaments to be $\\sim 1.45 \\times 10^{30}$ ergs, which is a small fraction of the total magnetic energy of $ \\sim 3 \\times 10^{32}$ ergs released in the flare \\citep{Schrijver08}. This mechanism of induced photospheric magnetic response to coronal magnetic restructuring appears energetically feasible to drive the lateral motion of the filaments. We call this first phase the ``rapid restructuring phase\" (RRP).  The slow global change in the displacement of the filaments during second phase can then be considered as a response of the gradual restructuring of the field on the much longer Alfven time-scale corresponding to the larger size of the CME flux rope system \\citep{Liu2008}. We call this phase the ``gradual restructuring phase\" (GRP). However, we need more observations to confirm these phases in other flares.  Vector magnetic field maps at a cadence similar to G-band images are not available from SOT/SP. However, future  instruments, like HMI aboard SDO mission,  will have high-resolution vector magnetic field maps at high-cadence and will shed more light on the flare related vector field changes associated with such lateral motions of small-scale features."
    },
    "0910/0910.4383_arXiv.txt": {
        "abstract": "We use our most recent training set for the {\\sc Rico} code to estimate the impact of recombination uncertainties on the posterior probability distributions which will be obtained from future CMB experiments, and in particular the {\\sc Planck} satellite. Using a Monte Carlo Markov Chain analysis to sample the posterior distribution of the cosmological parameters, we find that {\\sc Planck} will have biases of $-0.7$, $-0.3$ and $-0.4$ sigmas for $\\nS$, $\\omegab$ and $\\logAs$, respectively, in the minimal six parameter $\\Lambda$CDM model, if the description of the recombination history given by {\\sc Rico} is not used. The remaining parameters (e.g. $\\tau$ or $\\omegadm$) are not significantly affected. We also show, that the cosmology dependence of the corrections to the recombination history modeled with {\\sc Rico} has a negligible impact on the posterior distributions obtained for the case of the {\\sc Planck} satellite. In practice, this implies that the inclusion of additional corrections to existing recombination codes can be achieved using simple cosmology-independent `fudge functions'.  Finally, we also investigated the impact of some recent improvements in the treatment of hydrogen recombination which are still not included in the current version of our training set for {\\sc Rico}, by assuming that the cosmology dependence of those corrections can be neglected. In summary, with our current understanding of the complete recombination process, the expected biases in the cosmological parameters inferred from {\\sc Planck} might be as large as $-2.3$, $-1.7$ and $-1$ sigmas for $\\nS$, $\\omegab$ and $\\logAs$ respectively, if all those corrections are not taken into account. We note that although the list of physical processes that could be of importance for {\\sc Planck} seems to be nearly complete, still some effort has to be put in the validation of the results obtained by the different groups.  The new {\\sc Rico} training set as well as the fudge functions used for this paper are publicly availabe in the {\\sc Rico}-webpage. ",
        "introduction": "The cosmic microwave background (CMB) is nowadays an essential tool of theoretical and observational cosmology. Recent advances in the observations of the CMB angular fluctuations in temperature and polarization \\citep[e.g. ][]{WMAP5-basic,WMAP5-params} provide a detailed description of the global properties of the Universe, and the cosmological parameters are currently known with accuracies of the order of few percent in many cases. The experimental prospect for the {\\sc Planck} satellite \\citep{Planck2006}, which was launched on May 14th 2009, is to achieve the most detailed picture of the CMB anisotropies down to angular scales of $\\ell \\sim 2500$ in temperature and $\\ell \\sim 1500$ in polarization. This data will achieve sub-percent precision in many cosmological parameters.  However, those high accuracies will rely on a highly precise description of the theoretical predictions for the different cosmological models. Currently, it is widely recognised that the major limiting factor in the accuracy of angular power spectrum calculations is the uncertainty in the ionization history of the Universe \\citep[see ][]{hu1995, Seljak2003}.  This has motivated several groups to re-examine the problem of cosmological recombination \\citep{Zeldovich1968, Peebles1968}, taking into account detailed corrections to the physical processes occurring during hydrogen \\citep[e.g.][]{Dubrovich2005, RHS2005, Chluba2006, Kholupenko2006, Rubino2006, Chluba2007, Chluba2007a, Karshenboim2008, Hirata2008, Chluba2009c, Chluba2009b, Chluba2009, Jentschura2009, Labzowsky2009, HirataForbes2009} and helium recombination \\citep[e.g][]{Kholupenko2007, Wong2007, Switzer2008a, Switzer2008b, HirataSwi2008, Rubino2008, Kholupenko2008, Chluba2009d}. Each one of the aforementioned corrections individually leads to changes in the ionization history at the level of $\\ga 0.1$\\%, in such a way that the corresponding overall uncertainty in the CMB angular power spectra exceeds the benchmark of $\\pm 3/\\ell$ at large $\\ell$ \\citep[for more details, see][hereafter FCRW09]{rico}, thus biasing any parameter constraints inferred by experiments like {\\sc Planck}, which will be cosmic variance limited up to very high multipoles.  The standard description of the recombination process is provided by the widely used {\\sc Recfast} code \\citep{Seager1999}, which uses effective three-level atoms, both for hydrogen and helium, with the inclusion of a conveniently chosen {\\it fudge factor} which artificially modifies the dynamics of the process to reproduce the results of a multilevel recombination code \\citep{Seager2000}.  The simultaneous evaluation of all the new effects discussed above make the numerical computations very time-consuming, as they currently require the solution of the full multilevel recombination code. Moreover, some of the key ingredients in the accurate evaluation of the recombination history (e.g. the problem of radiative transfer in hydrogen and the proper inclusion of two-photon processes) are solved using computationally demanding approaches, although in some cases semi-analytical approximations \\citep[see e.g.][]{Hirata2008} might open the possibility of a more efficient evaluation in the future.  In order to have an accurate and fast representation of the cosmological recombination history as a function of the cosmological parameters, two possible approaches have been considered in the literature. The first one consists of the inclusion of additional fudge factors to mimic the new physics, as recently done in \\cite{Wong2008} (see {\\sc Recfast} v1.4.2), where they include an additional fudge factor to modify the dynamics of helium recombination. The second approach is the so-called {\\sc Rico} code (FCRW09), which provides an accurate representation of the recombination history by using a regression scheme based on {\\sc Pico} \\citep{Fendt2007a,Fendt2007}. The {\\sc Rico} code smoothly interpolates the $X_{\\rm e}(z; \\vec{p})$ function on a set of pre-computed recombination histories for different cosmologies, where $z$ is the redshift and $\\vec{p}$ represents the set of cosmological parameters.  In this paper, we present the results for parameter estimations using {\\sc Rico} with the most recent training set presented in FCRW09. This permits us to accurately account for the full cosmological dependence of the corrections to the recombination history that were included in the multi-level recombination code which was used for the training of {\\sc Rico} (see Sect.~\\ref{sec:rico} for more details). With this tool, we have evaluated the impact of the corrections on the posterior probability distributions that are expected to be obtained for the {\\sc Planck} satellite, by performing a complete Monte Carlo Markov Chain analysis. The study of these posteriors have shown that the impact of the cosmology dependence is not very relevant for those processes included into the current {\\sc Rico} training set. Therefore, by assuming that the cosmology dependence of the correction in general can be neglected, we have also investigated the impact of recent improvements in the treatment of hydrogen recombination (see Sect.~\\ref{sec:Lya}). The basic conclusion is that, with our current understanding of the recombination process, the expected biases in the cosmological parameters inferred from {\\sc Planck} might be as large as 1.5-2.5 sigmas for some parameters as the baryon density or the primordial spectral index of scalar fluctuations, if all these corrections to the recombination history are neglected.  The paper is organized as follows. Sect.~\\ref{sec:physics} describes the current training set for {\\sc Rico}, and provides an updated list of physical processes during recombination which were not included in FCRW09.  Sect.~\\ref{sec:impact} presents the impact of the recombination uncertainties on cosmological parameter estimation, focusing on the case of {\\sc Planck} satellite. Sect.~\\ref{sec:additional} further extends this study to account for the remaining recombination uncertainties described in Sect.~\\ref{sec:physics}. Sect.~\\ref{sec:current} presents the analysis of present-day CMB experiments, for which the effect is shown to be negligible. Finally, the discussion and conclusions are presented in sections \\ref{sec:discussion} and \\ref{sec:conclusions}, respectively. ",
        "conclusions": "\\label{sec:conclusions}  In this paper, we have performed a MCMC analysis of the expected biases on the cosmological constraints to be derived from the upcoming {\\sc Planck} data, in the light of recent developments in the description of the standard cosmological recombination process. Our main conclusions are: \\begin{itemize} \\item An incomplete description of the cosmological recombination process leads to significant biases (of several sigmas) in some of the basic parameters to be constrained by {\\sc Planck} satellite (see Table~\\ref{tab:planck2}), and in general, by any future CMB experiment. However, these corrections have a minor impact for present-day CMB experiments; for instance, using WMAP5 data plus other cosmological datasets, we find a $\\sim -0.25$ and $\\sim -0.22$ sigma bias on $\\nS$ and $\\logAs$, respectively, while the rest of the parameters remain unchanged.  \\item Today, it seems that our understanding of cosmological recombination has reached the sub-percent level in $X_{\\rm e}$ at redshifts $500 \\la z \\la 1600$. However, it will be important to cross-validate all of the considered corrections in a detailed code comparison, which currently is under discussion among the different groups.  \\item Given the range of variation of the relevant cosmological parameters, it is possible to incorporate all the new recombination corrections by using (cosmology independent) fudge functions. Here we described one possibility which uses a simple correction factor to the results obtained with {\\sc Recfast} (see Sect.~\\ref{sec:corr_fac}). We provide the function $f(z)$ on the {\\sc Rico}-webpage \\footnote{http://cosmos.astro.uiuc.edu/rico}.  \\item The physics of helium recombination already seems to be captured at a sufficient level of precision, when including the acceleration caused by the hydrogen continuum opacity and the $2^3\\rm P_1$-$1^1\\rm S_0$ intercombination line. The biases caused by neglecting only these corrections are -0.8 and -0.4 sigmas, for $\\nS$ and $\\omegab$, respectively.  \\item When allowing for more non-standard additions to the recombination model (e.g. related to annihilating dark matter), the biases introduced by an inaccurate recombination model could lead to spurious detections or additional confusion (see Sect.~\\ref{sec:nonstandard}).  \\end{itemize}"
    },
    "0910/0910.3330_arXiv.txt": {
        "abstract": "\\noindent The Carnegie Supernova Project is a five-year survey being carried out at the Las Campanas Observatory to obtain high-quality light curves of $\\sim$100 low-redshift Type~Ia supernovae in a well-defined photometric system.  Here we present the first release of photometric data that contains the optical light curves of 35 Type~Ia supernovae, and near-infrared light curves for a subset of 25 events. The data comprise 5559 optical ($ugriBV$) and 1043 near-infrared ($YJHK_s$) data points in the natural system of the Swope telescope. Twenty-eight supernovae have pre-maximum data, and for 15 of these, the observations begin at least 5 days before $B$ maximum.  This is one of the most accurate datasets of low-redshift Type~Ia supernovae published to date. When completed, the CSP dataset will constitute a fundamental reference for precise determinations of cosmological parameters, and serve as a rich resource for comparison with models of Type~Ia supernovae. ",
        "introduction": "\\label{sec:intro}  The observation that the expansion rate of the Universe is currently accelerating is arguably one of the most profound discoveries in modern astrophysics.  The first direct evidence of this phenomenon was provided a decade ago by the Hubble diagram of high-redshift Type~Ia supernovae \\citep[SNe~Ia;][]{riess98,perlmutter99}. Since then, a wealth of data obtained from surveys of nearby and distant SNe~Ia has confirmed this conclusion, as have other methods such as the X-ray cluster distances \\citep[e.g.,][]{allen07} and the late-time integrated Sachs-Wolfe effect using cosmic microwave background radiation observations \\citep[e.g.,][]{giannantonio08}. These findings suggest that a new form of energy permeates the Universe, or that the theory of General Relativity breaks down on cosmological scales.  Today, the major challenge is to determine the nature of this mysterious energy (commonly referred to as ``dark energy'') by measuring its equation-of-state parameter, $w = P/(\\rho c^2)$, and the corresponding time derivative $\\dot w$.  SNe~Ia are playing an essential role in the endeavor to measure $w$. Both the SN Legacy Survey \\citep{astier06} and the ESSENCE project \\citep{wood-vasey07} have recently provided independent constraints on $w$ that favor a cosmological constant ($w = -1$).  The determination of $\\dot w$, however, will require new and extensive samples of both low- and high-redshift SNe~Ia with systematic errors below 1--2\\%.  This will require the construction of a database of low-redshift light curves with excellent photometric precision and temporal coverage   The pioneering work of producing such a low-redshift sample of SNe~Ia was carried out by the Cal{\\' a}n/Tololo Survey, which published $BVRI$ light curves of 29 events \\citep{hamuy96}. This is the local sample that was used by \\citet{riess98} and \\citet{perlmutter99} to detect the accelerating expansion of the Universe.  In 1999, the CfA Supernova Group released a set of 22 $BVRI$ light curves \\citep{riess99}.  Subsequently, this group has published $UBVRI$ light curves of 44 events \\citep{jha06}, and $UBVRIr'i'$ light curves of another 185 \\citep{hicken09}. Other significant samples of nearby SNe~Ia are being produced by the Lick Observatory Supernova Search \\citep[LOSS;][]{filippenko01,filippenko05,filippenko09} and the Nearby Supernova Factory \\citep{aldering02}. In addition, the SDSS-II Supernova Survey \\citep{frieman08} has obtained $u'g'r'i'z'$ light curves of a sample of $\\sim$500 spectroscopically confirmed SNe~Ia in the redshift range of $0.05 < z < 0.35$.  In September 2004, the Carnegie Supernova Project (hereafter CSP) began a five-year program to obtain densely sampled optical {\\em and} near-infrared light curves of $\\sim$100 nearby SNe~Ia.  The overriding goal has been to obtain the highest possible precision in a stable, well-understood photometric system.  An additional objective is to use the broad wavelength leverage afforded by these observations to set stringent constraints on the properties of host-galaxy extinction \\citep[see][]{folatelli09}.  In this paper, we present final photometry of the first 35 SNe~Ia observed in the years 2004--2006. This dataset consists of optical ($u'g'r'i'BV$) light curves of all 35 SNe, and near-infrared ($YJHK_s$) light curves for a subset of 25.  In total, 5559 optical and 1043 near-infrared measurements were obtained at a typical precision of 0.01--0.03~mag, making this one of the most homogeneous and accurate sets of nearby SN~Ia light curves yet obtained.  A distinguishing aspect of the CSP is the quantity and quality of the near-infrared photometry, which is matched only by the recently published PAIRITEL results \\citep{wood08}.  Previously, \\citet[][hereafter H06]{hamuy06} gave a detailed description of the CSP observing methodology and data reduction procedures, and showed several examples of light curves obtained during the first campaign.  In a second paper, \\citet{folatelli06} reported on the unique Type Ib SN~2005bf, while the third CSP paper presented a detailed analysis of the peculiar Type~Ia SN 2005hk \\citep{phillips07}.  A fourth paper on the distance to the Antennae Galaxies (NGC~4038/39) based on the Type~Ia SN~2007sr was presented by \\citet{schweizer08}.  CSP photometry has also been included in case studies of the subluminous Type~Ia SN~2005bl \\citep{taubenberger08} and the Type~Ia/IIn SN~2005gj \\citep{prieto07}. Most recently, \\citet{stritzinger09} published a comprehensive study on the Type~Ib SN~2007Y.  The organization of this paper is as follows. Section~\\ref{sec:obs} describes the observing and basic data reduction procedures, Section~\\ref{sec:seq} reviews the establishment of the local photometric sequences, Section~\\ref{sec:phot} gives details of the measurement and calibration of the SN photometry, and Section~\\ref{sec:results} presents the final light curves.  In two Appendices, we present our most up-to-date information on the sensitivity functions of our $u'g'r'i'BV$ bandpasses, and near-infrared photometry of Feige~16 obtained over the course of our observing campaigns which we use to check the validity of our $Y$-band calibration.  Accompanying this paper are two other papers. The first gives an analysis of the data presented here \\citep{folatelli09}. The second combines the data of this article with that of 35 high-redshift SNe~Ia, observed in the near-infrared with the Magellan Baade 6.5-m telescope, in order to construct the first rest-frame $i$-band Hubble diagram out to $z \\approx 0.7$ \\citep{freedman09}. ",
        "conclusions": "\\label{sec:conc}  In this paper, we have presented the first data release of low-redshift SNe~Ia observed by the CSP.  The combination of pre-maximum observations, extended coverage, dense sampling, and high signal-to-noise ratio data produces a dataset of unprecedented quality.  These results will make possible detailed studies of the photometric properties of SNe~Ia and serve as a rich resource for comparison with theoretical models.  When completed, the CSP dataset will constitute a fundamental reference for precise determinations of cosmological parameters such as our own program to measure the $i'$-band Hubble diagram of SNe~Ia to $z \\approx 0.7$ \\citep{freedman09}. The inclusion of the near-infrared data is an important addition since SNe~Ia are nearly perfect standard candles at these wavelengths \\citep{krisciunas04}.  In addition, the broad wavelength range covered by the optical and near-infrared observations provides strong leverage for estimating dust extinction in the SN~Ia host galaxies \\citep{folatelli09}."
    },
    "0910/0910.4726_arXiv.txt": {
        "abstract": "{ In this paper, the stability of a dynamically condensing radiative gas layer is investigated by linear analysis. Our own time-dependent, self-similar solutions describing a dynamical condensing radiative gas layer are used as an unperturbed state. We consider perturbations that are both perpendicular and parallel to the direction of condensation. The transverse wave number of the perturbation is defined by $k$. For $k=0$, it is found that the condensing gas layer is unstable. However, the growth rate is too low to become nonlinear during dynamical condensation. For $k\\ne0$, in general, perturbation equations for constant wave number cannot be reduced to an eigenvalue problem due to the unsteady unperturbed state. Therefore, direct numerical integration of the perturbation equations is performed. For comparison, an eigenvalue problem neglecting the time evolution of the unperturbed state is also solved and both results agree well. The gas layer is unstable for all wave numbers, and the growth rate depends a little on wave number. The behaviour of the perturbation is specified by $kL_\\mathrm{cool}$ at the centre, where the cooling length, $L_\\mathrm{cool}$, represents the length that a sound wave can travel during the cooling time. For $kL_\\mathrm{cool}\\gg1$, the perturbation grows isobarically. For $kL_\\mathrm{cool}\\ll1$, the perturbation grows because each part has a different collapse time without interaction. Since the growth rate is sufficiently high, it is not long before the perturbations become nonlinear during the dynamical condensation. Therefore, according to the linear analysis, the cooling layer is expected to split into fragments with various scales. } ",
        "introduction": "In the interstellar medium (ISM), it is well known that a clumpy low-temperature phase (cold neutral medium, or CNM) and a diffuse high-temperature phase (warm neutral medium, or WNM) can coexist in pressure equilibrium as a result of the balance of radiative cooling and heating due to external radiation fields and cosmic rays \\citep{FGH69,W95,W03}. These two phases are thermally stable. On the other hand, gas is thermally unstable in the temperature range between two stable phases, that is, in the range 300 K $<T<$ 6000 K. The unstable gas spontaneously turns into the mixture of stable CNM and WNM by thermal instability (TI). This instability is expected to play an important role in the structure formation and the dynamics of the ISM, and especially, in the molecular cloud formation and origin of turbulence.  The basic properties of TI was investigated by \\citet{F65}, who performed linear analysis of an uniform gas in thermally equilibrium. He derived a criterion for TI. Focusing on one fluid element, \\citet{B86} generalized the Field criterion when the cooling rate is not equal to the heating rate. The effect of magnetic field on TI has been investigated by \\citet{F65} and \\citet{HP06}, and other authors.  Recently, many authors have used multi-dimensional numerical simulations to study the turbulent CNM formation driven by TI. \\citet{KI02} suggest that the turbulent CNM formation is induced by TI in a shock-compressed region. Analogous processes have been studied by many authors for a colliding flow of the WNM \\citep{AH05, HA07, H06,V07}, and using two-fluid MHD simulation \\citep{II08}. The unbalance between cooling and heating rates causes the shock-compressed gas layer to cool and to condense. During the cooling, these numerical simulations shows that the runaway cooling layer quickly fragments into many CNM clumps whose velocity dispersion is equal to a fraction of the sound speed of WNM, where CNM clumps and WNM are tightly interwoven. This complex structure is regarded as produced by TI and possibly by some other hydrodynamical instabilities, such as the nonlinear thin shell instability \\citep{V94}, the Kelven-Helmholtz instability, and by corrugation instability of the phase transition layers between CNM and WNM \\citep{IIK06}.  A fluid element that is compressed by a shock wave tends to be a layer rather than a sphere because it is only compressed in the direction perpendicular to the shock front. Once the fluid element enters the thermally unstable regime, the layer cools in a runaway fashion. In this paper, we focus on the fragmentation of the runaway cooling layer. In previous studies, A detailed physical mechanism of the fragmentation of the runaway cooling layer remains poorly understood even in linear regime. The main reason is that it is difficult to select the unperturbed state since the cooling layer evolves temporarily and spatially. Therefore, in previous works, the unperturbed states were limited to spatially uniform gas that cools isochorically \\citep{SMS72,BL00} or isobarically \\citep{KI00}.  \\citet{IT08} (hereafter IT08) have recently found a family of self-similar (S-S) solutions describing the dynamical condensation of a radiative gas layer where the cooling rate dominates the heating rate. This S-S solution assumed that the net cooling rate is a power-law function and that the heating rate is explicitly neglected. Although it is still too ideal, they are expected to be a good nonlinear one-dimensional model at least in the phase during the transition from WNM to CNM. In this paper, we adopt the S-S solutions as a more realistic unsteady unperturbed state than those in previous works. We perform linear analysis of the S-S solutions against fluctuations perpendicular, as well as parallel to, the direction of the condensation. By performing the linear analysis, we will have some useful insights when and how the cooling layer fragments. Since we focus on the above S-S unperturbed state, the nonlinear thin shell instability, Kelvin-Helmholtz instability, and the corrugation instability are beyond the scope of this paper.  In Sect. \\ref{formulation}, we formulate basic equations using a zooming coordinate. Perturbation equations are derived for the linear analysis, with a brief review of the S-S solutions. In Sect. \\ref{keq0}, we investigate the stability of the S-S solutions taking only those fluctuations into account that are parallel to the direction of the condensation. In Sect. \\ref{sec kne0}, we consider the fluctuations that are both perpendicular and parallel to the direction of the condensation. In Sect. \\ref{discuss}, we discuss the astrophysical implication of the linear analysis and effects of the thermal conduction. Our study is summarized in Sect. \\ref{summary}. ",
        "conclusions": "\\label{discuss} \\subsection{Astrophysical implication}\\label{astro impli} In this section, we estimate the fragmentation time of the cooling layer formed by a collision of WNM. We consider a head-on collision between two thermally equilibrium gases in which density and pressure are $\\rho_\\mathrm{WNM}=0.57m_\\mathrm{H}\\;\\mathrm{cm}^{-3}$ and $P_\\mathrm{WNM}=3.5\\times10^3k_\\mathrm{B}\\;\\mathrm{dyne\\;cm^{-2}}$. Two clouds collide along the $x$-axis at $t=0$ and $x=0$ with velocity $20\\;\\mathrm{km\\;s^{-1}}$, i.e., the mach number is 2.17. The cooling function in this temperature region fitted by \\citet{KI02} as follows: \\begin{equation} \\rho{\\cal L} = n (-\\Gamma + n\\Lambda)\\;\\;\\mathrm{erg\\;cm^{-3}\\;s^{-1}}, \\end{equation} where \\begin{equation} \\Gamma = 2\\times10^{-26}, \\end{equation} \\begin{equation} \\frac{\\Lambda}{\\Gamma} = 1.0\\times10^7 \\exp\\left( -\\frac{118400}{T+1000} \\right) + 1.4\\times10^{-2}\\sqrt{T}\\exp\\left( -\\frac{92}{T} \\right), \\end{equation} where $n$ is the number density of the gas. We perform numerical calculation using the one-dimensional Lagrangian Godunov scheme \\citep{L97} between $x=0$ and $x=L_x = 3.26$pc. At $x=0$ and $x=L_x$, we adopt the reflecting and free boundary conditions, respectively. As the initial condition, we add the following density fluctuation: \\begin{equation} \\rho(t=0,x)=\\frac{\\rho_\\mathrm{WNM}} {\\displaystyle 1+\\frac{A}{i_\\mathrm{max}}\\sum_{i=0}^{i_\\mathrm{max}} \\sin\\left( \\frac{2\\pi }{L_x} +\\theta_i\\right)    }, \\end{equation} where we set $i_\\mathrm{max}=10,\\;A=0.5$, and the phase, $\\theta_i$, is given by random number between 0 and $2\\pi$. The maximum and minimum number density at the initial state are 0.72 cm$^{-3}$ and 0.44 cm$^{-3}$, respectively.  \\begin{figure}[h] \\begin{center} \\includegraphics{12806fg7.eps} \\end{center} \\caption{Evolutionary track of the gas at the centre after the head-on collision (the thick solid line). The ordinate and abscissa axes indicate the pressure and the number density, respectively. The thin solid line indicates the thermally equilibrium curve. The filled circle represents the gas at the preshock region. } \\label{Pden evo} \\end{figure}  \\begin{figure}[h] \\begin{center} \\includegraphics{12806fg8.eps} \\end{center} \\caption{ Time evolution of $\\alpha_\\mathrm{num}$ (the solid line) and the cooling length (the dotted line), which are evaluated at the centre . The dashed line indicates the time evolution of the critical wave length defined in Sect. \\ref{thermal conduction}. } \\label{alpha evo} \\end{figure}  \\begin{figure*}[t] \\begin{center} \\includegraphics{12806fg9.eps} \\end{center} \\caption{ Time evolution of (a)the number density and (b)the temperature, respectively. The rescaled (c) number density, $nt_*^{\\eta/(2-\\alpha)}$ and (d) temperature, $T t_* ^{(1-\\alpha)(1-\\eta)/\\{\\eta(2-\\alpha)\\}}$ as a function of the rescaled coordinate, $xt_*^{-1/(1-\\omega)}$, where $t_\\mathrm{c}$ and $\\eta$ are set to 0.915 Myr and 0.98, respectively. The thick lines in (c) and (d) indicate the corresponding S-S solution. } \\label{inter} \\end{figure*}  Figure \\ref{Pden evo} shows the evolutionary track of the gas at the centre, $x=0$. Due to the shock compression, the temperature of gas increases suddenly, and the postshock region enters the thermally unstable phase from the initial stable state. The TI leads to the cooling layer condensing in a runaway fashion. The cooling length in Fig. \\ref{alpha evo} rapidly decreases with time until it reaches minimum value $1.5\\times10^{-3}$ pc at $t=0.917$ Myr. During runaway cooling, the gas condenses isobarically until it reaches the stable CNM phase, although some pressure oscillation is seen in Fig. \\ref{Pden evo}. We focus on the property of this runaway condensing layer. First we evaluate $\\alpha_\\mathrm{num}(t)$, which is defined by \\begin{equation} \\alpha_\\mathrm{num}=\\left( \\frac{ \\partial (\\rho{\\cal L})}{\\partial T} \\right) _\\mathrm{\\rho}\\Biggr|_{x=0} \\end{equation} at each instant of time. Figure \\ref{alpha evo} shows the time evolution of $\\alpha_\\mathrm{num}$. During $0.4<t/\\mathrm{Myr}<0.9$, it is seen that $\\alpha_\\mathrm{num}\\sim0.61$ is approximately constant. Therefore, in this period, the flow is expected to be approximated well by the S-S solutions.  Figures \\ref{inter} show the time evolutions of the distributions of (a) the number density and (b) the temperature at $t=$ 0.72, 0.80, 0.88, and 0.90 Myr, respectively. The rescaled number density $nt_*^{\\eta/(2-\\alpha)}$ and temperature, $T t_*^{\\left\\{ \\eta - (2-\\alpha) \\right\\}/\\{(2-\\alpha)(1-\\alpha)\\}}$ are shown in Figs. \\ref{inter}c and \\ref{inter}d as a function of rescaled coordinate, $xt_*^{-1/(1-\\omega)}$, where $t_\\mathrm{c}$, $\\eta$, and $\\alpha$ are set to 0.915 Myr, 0.98, and 0.61, respectively. From Figs. \\ref{inter}c and \\ref{inter}d, it is clearly seen that the S-S solution $(\\alpha=0.61, \\eta=0.98)$ describes the results of the numerical calculation well. There are two reasons they are described by the S-S solution well. One is that $\\alpha_\\mathrm{num}$ is almost constant, and the cooling function is approximated well by Eq. (\\ref{cooling func}) during the runaway condensation. The other reason relates to the stability of S-S solutions. The one-dimensional calculation solves the evolution only in the direction parallel to the condensation. From the linear analysis in Sect. \\ref{keq0}, such a perturbation with $k=0$ cannot grow enough during the runaway condensation. That is why the S-S solution is realized even though the density is initially fluctuated.  However, the result is expected to be qualitatively different if the perturbation with $k\\ne0$ exists. The perturbation with $k\\ne0$ on the isobarically runaway condensing layer grows as $\\propto t_*^{-1}$ (see Sect. \\ref{sec kne0}). The growth rate is independent of wave numbers. From Figs. \\ref{alpha evo} and \\ref{inter}, the gas evolves obeying the S-S solution between $t=$ 0.4 and 0.9 Myr. If the layer has a perturbation ($k\\ne0$) with amplitude $\\delta \\rho_0$ at $t=0.4\\;\\mathrm{Myr}$, the perturbation grows as \\begin{equation} \\frac{\\delta \\rho(t)}{\\delta \\rho_0}=\\frac{1-0.4 \\mathrm{Myr}/t_\\mathrm{c}} {1 - t/t_\\mathrm{c}}. \\end{equation} The epochs when the perturbation grows by a factor of 10 and 100 are $t=$ 0.86 and 0.91 Myr, respectively. At 0.86 Myr, the central number density is still as low as 38 cm$^{-3}$. Therefore, the condensing layer is expected to fragment quickly, and CNM clumps will form.  \\cite{II08} investigated TI in a shock compressed region formed by WNM-WNM collision using two-dimensional, two-fluid magnetohydrodynamical simulation. The initial condition there is the same as that in our one-dimensional simulation except for the dimension and the way to add density fluctuation in the WNM. Our linear analysis is applicable to the gas evolution before CNM clouds form in their unmagnetized case. In the initial stage of their simulation, in the postshock region, thermally unstable gas initially condenses in a layer structure. Around $t=0.5$ Myr when the region of the highest density reaches CNM stable state, both small CNM cloudlets and filamentary CNM clouds are generated \\citep{I09}. According to our linear analysis, an isobarically condensing gas layer is expected to split into fragments that have a variety of lengths in the transverse direction. Therefore, our linear analysis agrees with their simulation qualitatively.  After the first CNM clouds formation, the mass of these CNM clouds continues to increase by the accretion of unstable gas by the shocks. Moreover, some CNM clouds coalesce into larger clouds and others fragment into smaller ones by the turbulent flow \\citep{I09}. Therefore, the size and mass of CNM clouds varies with time. To understand this phase, which is beyond the scope of this paper, a statistical approach is probably needed. A statistical theory has been proposed by \\cite{HA07} using Press-Schechter formalism \\citep{PS74}, assuming that the CNM clumps are generated from density fluctuations within WNM whose power spectrum is assumed to be Kolmogorov.  Comparing timescales, \\cite{HHB08} discuss the effect of TI, the nonlinear thin-shell instability, Kelvin-Helmholtz instability, and gravitational instability in various densities, temperatures, and scales of fluctuations. In their paper, it is assumed that the gas is unstable only if the scale of the fluctuation is smaller than the sound crossing scale. However, from our linear analysis, perturbations can grow regardless of their scales as long as they are flat rather than spherical.  \\subsection{The growth rate for $1<\\alpha<2$} Although the linear analysis on the S-S cooling layer is limited for $\\alpha<1$, the thermal stability of the gas for $\\alpha>1$ is also roughly understood from Balbus's criterion. For $1<\\alpha<2$, the gas  is isobarically unstable, but it is isochorically stable. For $\\alpha>2$, the gas  is thermally stable. In this section, we investigate the stability of the gas for $1<\\alpha<2$ during cooling within the large and small scale limits.  \\subsubsection{The isobaric mode} For the case with small wave length, perturbation is expected to grow isobarically. By comparison of our results with previous studies in the literature, it is found that the growth rate in the isobaric mode is independent of the global structure of the unperturbed state. Therefore, from local arguments, the growth rate in the isobaric mode of the gas with $1<\\alpha<2$ can also be estimated.  As an unperturbed state, we adopt a cooling gas element whose scale is assumed to be much smaller than the cooling length. In this case, the element cools isobarically. From Eq. (\\ref{eoe}), the time evolution of the unperturbed gas is given by \\begin{equation} \\rho(t)=\\rho_i\\left( 1-\\frac{t}{t_\\mathrm{cool}'} \\right)^{-1/(2-\\alpha)}, \\frac{1}{t_\\mathrm{cool}'} = (2-\\alpha)\\gamma^{\\alpha-1}(\\gamma-1) P_i^{\\alpha-1}\\rho_i^{2-\\alpha}, \\label{a ge 1} \\end{equation} where $\\rho_i$ and $P_i$ represent the initial density and pressure, respectively. In the above unperturbed state, we consider the following isobaric perturbation: \\begin{equation} \\rho = \\rho_0(t) +\\delta \\rho(t), \\end{equation} and \\begin{equation} P = P_0, \\end{equation} where subscript ``0'' indicates the unperturbed state, and $\\delta \\rho$ is the density perturbation. Linearizing Eq. (\\ref{eoe}), one obtains \\begin{equation} \\frac{\\mathrm{d} }{\\mathrm{d}t}\\left( \\frac{\\delta\\rho}{\\rho_0} \\right) = (2-\\alpha)\\gamma\\gamma^{\\alpha-1} \\rho_0^{2-\\alpha}P_0^{\\alpha-1}\\frac{\\delta \\rho}{\\rho_0}. \\label{per a ge 1} \\end{equation} Using Eq. (\\ref{a ge 1}), Eq. (\\ref{per a ge 1}) is rewritten as \\begin{equation} \\frac{\\mathrm{d} }{\\mathrm{d}t}\\left( \\frac{\\delta\\rho}{\\rho_0} \\right) = \\frac{1}{t_\\mathrm{cool}'} \\left( 1-\\frac{t}{t_\\mathrm{cool}'} \\right)^{-1} \\frac{\\delta \\rho}{\\rho_0}. \\label{per a ge 1 2} \\end{equation} Equation (\\ref{per a ge 1 2}) is easily integrated to give \\begin{equation} \\frac{\\delta \\rho}{\\rho_0} \\propto \\left( 1-\\frac{t}{t_\\mathrm{cool}'} \\right)^{-1}. \\label{a ge 1 growth rate} \\end{equation} Comparing Eq. (\\ref{a ge 1 growth rate}) with Eq. (\\ref{a ge 1}), one can see that the perturbation grows more slowly than the unperturbed state for $1<\\alpha<2$. Therefore, the gas is expected to be difficult to fragment during runaway cooling if $1<\\alpha<2$.  \\subsubsection{The noninteractive mode} A cooling layer that evolves isobarically is considered. The time evolution of the central density is the same as Eq. (\\ref{a ge 1}). When the scale of perturbation perpendicular to the condensation is too large to interact with other regions, each region evolves independently. Here, we focus on the time evolution of density perturbation at the centre ($x=0$). Initial fluctuation of the central density, $\\delta \\rho_i$, creates the fluctuation of the cooling time, $\\Delta t$. The relative amplitude of density perturbation at $x=0$ is given by \\begin{equation} \\frac{\\delta \\rho}{\\rho_0} = \\frac{1}{\\rho_0}\\left( \\rho_i + \\delta \\rho_i \\right) \\left( 1-\\frac{t}{t_\\mathrm{cool}' - \\Delta t} \\right)^{-1/(2-\\alpha)}-1. \\label{non int a ge 1} \\end{equation} Linearizing Eq. (\\ref{non int a ge 1}) with omitting terms that do not grow, we have \\begin{equation} \\frac{\\delta \\rho}{\\rho_0} = \\frac{1}{2-\\alpha}\\frac{\\Delta t}{t_\\mathrm{cool}'} \\left( 1-\\frac{t}{t_\\mathrm{cool}'} \\right)^{-1}. \\label{non in a ge 1 1} \\end{equation} Comparing Eq. (\\ref{non in a ge 1 1}) with Eq. (\\ref{a ge 1}), we can see that the perturbation grows more slowly than the unperturbed state for $1<\\alpha<2$. Therefore, the gas is expected to be difficult to fragment for $1<\\alpha<2$ for the large-scale perturbation, as well as the small scale.  \\subsection{Effects of thermal conduction}\\label{thermal conduction} In this paper, the thermal conduction is neglected for simplicity. However, for large wave number, the thermal conduction is expected to stabilize TI in the cooling layer \\citep{F65}. Therefore, there is a critical wave number, $k_\\mathrm{crit}$, such that perturbation with a larger wave number is stabilized by the thermal conduction.  First, we evaluate $k_\\mathrm{crit}$ using an order estimation. Using the characteristic time scale of the thermal conduction, $t_\\mathrm{diff}$, the diffusion equation is given by \\begin{equation} \\frac{1}{t_\\mathrm{diff}} \\left( \\frac{P_{00}}{\\gamma-1} \\right) \\sim k^2K(T_{00}) T_{00}, \\label{diffusion eq} \\end{equation} where $K$ is the thermal conduction coefficient. From Eq. (\\ref{diffusion eq}), the diffusion timescale is given by \\begin{equation} t_\\mathrm{diff} \\simeq \\frac{P_{00}}{(\\gamma-1)K(T_{00})T_{00} } k^{-2}. \\label{diffusion} \\end{equation} From Eq. (\\ref{diffusion}), one can see that the diffusion timescale is small for large wave number. If $t_\\mathrm{diff}<t_\\mathrm{cool}$, the thermal conduction is expected to stabilize TI. Therefore, $k_\\mathrm{crit}$ can be derived on the condition $t_\\mathrm{diff}\\sim t_\\mathrm{cool}$ as \\begin{equation} k_\\mathrm{crit} = \\sqrt{\\frac{k_\\mathrm{K}}{L_\\mathrm{cool}}},\\;\\;\\mathrm{where}\\; k_\\mathrm{K} = \\frac{P_{00}c_{00}}{(\\gamma-1)KT_{00}}. \\label{critical wave} \\end{equation} \\cite{F65} derived similar critical wave number to Eq. (\\ref{critical wave}). Since the unperturbed state is time dependent, $k_\\mathrm{crit}$ also evolves with time. Detailed evolution of $k_\\mathrm{crit}$ depends on $K$. In $T<6000$, we adopt $K=2.5\\times10^3\\sqrt{T}$ ergs cm$^{-1}$ K$^{-1}$ s$^{-1}$ \\citep{P53}. In this case, from Eq. (\\ref{critical wave}), the time evolution of $k_\\mathrm{crit}$ can be derived analytically as \\begin{eqnarray} \\frac{\\mathrm{d} \\ln k_\\mathrm{crit}}{\\mathrm{d}\\ln t_*} &=& \\frac{(2\\alpha-1)\\eta - (2-\\alpha)(3-2\\alpha)}{4(2-\\alpha)(1-\\alpha)}\\nonumber\\\\ &=&\\left\\{\\begin{array}{cc} \\displaystyle - \\frac{7-2\\alpha}{4(2-\\alpha)}<0 & \\mathrm{for}\\;\\eta=1\\\\ \\displaystyle - \\frac{3-2\\alpha}{4(1-\\alpha)}<0 & \\mathrm{for}\\;\\eta=0\\\\ \\end{array}\\right.. \\label{thermal cond} \\end{eqnarray} For $\\eta=1$ and $\\eta=0$, it is found that Eq. (\\ref{thermal cond}) is negative for $\\alpha<1$. Since Eq. (\\ref{thermal cond}) is the linear function for $\\eta$, $\\mathrm{d}\\ln k_\\mathrm{crit}/\\mathrm{d}\\ln t_*$ is negative for all $\\eta$. Therefore, $k_\\mathrm{crit}$ increases with time. This means that an initially stable perturbation with wave number ($k>k_\\mathrm{crit}$) becomes unstable at a certain epoch when $k_\\mathrm{cirt}$ catches up with $k$.  Figure \\ref{alpha evo} shows the time evolution of the critical wave length, $\\lambda_\\mathrm{crit}=2\\pi/k_\\mathrm{crit}$. In Sect. \\ref{astro impli}, the time evolution of the condensing layer can be described by the S-S solution with ($\\alpha_\\mathrm{num}=0.61$, $\\eta=0.98$). In Fig. \\ref{alpha evo}, we see that $\\lambda_\\mathrm{crit}$ decreases with time during the runaway condensation. Figure \\ref{alpha evo} also shows that $\\lambda_\\mathrm{crit}<L_\\mathrm{cool}$, or $\\kappa_\\mathrm{crit}= k_\\mathrm{crit}L_\\mathrm{cool}>1$. This means that the effect of thermal conduction on the dispersion relation (Fig. \\ref{kneq0 grow}) always appears in the isobaric regime during the runaway condensation."
    },
    "0910/0910.0333_arXiv.txt": {
        "abstract": "{ Over almost all of minimal supergravity (mSUGRA or CMSSM) model parameter space, there is a large overabundance of neutralino cold dark matter (CDM). We find that the allowed regions of mSUGRA parameter space which match the measured abundance of CDM in the universe are highly fine-tuned. If instead we invoke the Peccei-Quinn-Weinberg-Wilczek solution to the strong $CP$ problem, then the SUSY CDM may consist of an axion/axino admixture with an axino mass of order the MeV scale, and where mixed axion/axino or mainly axion CDM seems preferred. In this case, fine-tuning of the relic density is typically much lower, showing that axion/axino CDM ($a\\tilde{a}$CDM) is to be preferred in the paradigm model for SUSY phenomenology. For mSUGRA with $a\\tilde{a}$CDM, quite different regions of parameter space are now DM-favored as compared to the case of neutralino DM. Thus, rather different SUSY signatures are expected at the LHC in the case of mSUGRA with $a\\tilde{a}$CDM, as compared to mSUGRA with neutralino CDM. } ",
        "introduction": "\\label{sec:intro}  A wide array of astrophysical data point to us living in a universe comprised of $4\\%$ baryons, $\\sim 25\\%$ cold dark matter (CDM) and $\\sim 70\\%$ dark energy. In fact, the cosmic abundance of CDM has been recently measured to high precision by the WMAP collaboration\\cite{wmap5}, which finds \\be \\Omega_{CDM}h^2=0.110\\pm 0.006 , \\ee where $\\Omega=\\rho/\\rho_c$ is the dark matter density relative to the closure density, and $h$ is the scaled Hubble constant. No particle present in the Standard Model (SM) of particle physics has the correct properties to constitue the CDM, so some form of new physics is needed. It is compelling, however, that candidate CDM particles do emerge naturally from two theories which provide solutions to longstanding problems in particle physics.  The first problem-- known as the gauge hierarchy problem-- arises due to quadratic divergences in the scalar sector of the SM. These divergences lead to scalar masses blowing up to the highest scale in the theory ({\\it e.g.} in grand unified theories (GUTS), the GUT scale $M_{GUT}\\simeq 2\\times 10^{16}$ GeV), unless an enormous fine-tuning of parameters is invoked. One solution to the gauge hierarchy problem occurs by introducing supersymmetry (SUSY) into the theory. The inclusion of softly broken SUSY leads to a cancellation of quadratic divergences between fermion and boson loops, so that only log divergences remain. The log divergence is soft enough that vastly different scales remain stable within a single effective theory. In SUSY theories, the lightest neutralino emerges as an excellent WIMP CDM candidate. Gravity-mediated SUSY breaking models (supergravity, or SUGRA) contain gravitinos with weak-scale masses. SUGRA models experience tension due to possible overproduction of gravitinos in the early universe, leading to an overabundance of CDM. In addition, gravitinos usually decay during or after Big Bang nucleosynthesis (BBN), and their energetic decay products may disrupt the successful calculations of light element abundances, which otherwise maintain good agreement with observation. This tension in SUGRA models is known as the {\\it gravitino problem}.  The second problem is the strong $CP$ problem\\cite{kcreview}. An elegant solution to the strong $CP$ problem was proposed by Peccei and Quinn (PQ) many years ago\\cite{pq}. The PQ solution automatically predicts the existence of a new particle (WW)\\cite{ww}: the axion $a$. While the original PQWW axion was soon ruled out, models of a nearly ``invisible axion'' were developed in which the PQ symmetry breaking scale was moved up to energies of order $f_a\\sim 10^{9}-10^{12}$ GeV\\cite{ksvz,dfsz}. The axion also turns out to be an excellent candidate particle for CDM in the universe\\cite{absik}.  Of course, it is highly desirable to simultaneously account for {\\it both} the strong $CP$ problem and the gauge hierarchy problem. In this case, it is useful to invoke supersymmetric models which include the PQWW solution to the strong $CP$ problem\\cite{nillesraby}. In a SUSY context, the axion field is just one element of an {\\it axion supermultiplet}. The axion supermultiplet contains a complex scalar field, whose real part is the $R$-parity even saxion field $s(x)$, and whose imaginary part is the axion field $a(x)$. The supermultiplet also contains an $R$-parity odd spin-$\\frac{1}{2}$ Majorana field, the axino $\\ta$\\cite{steffen_rev}. The saxion, while being an $R$-parity even field, nonethless receives a SUSY breaking mass likely of order the weak scale. The axion mass is constrained by cosmology and astrophysics to lie in a favored range $10^{-2}$ eV$\\agt m_a\\agt 10^{-5}$ eV. The axino mass is very model dependent\\cite{axmass,rtw,cl,ckkr}, depending heavily on the exact form of the superpotential and the mechanism for SUSY breaking. In supergravity models, it may be of order the gravitino mass $m_{3/2}\\sim$ TeV, or as low as $m_{3/2}^2/f_a\\sim $keV. Conditions for realizing these extremes are addressed in \\cite{cl}. Here, we will try to avoid explicit model-dependence, and adopt $m_{\\ta}$ as lying within the general range of keV-GeV, as in numerous previous works\\cite{rtw,ckkr,fstw,cmssm,axdm}. An axino in this mass range would likely serve as the lightest SUSY particle (LSP), and is also a good candidate particle for cold dark matter\\cite{rtw,ckkr}.  In a previous paper\\cite{axdm}, we investigated supersymmetric models wherein the PQ solution to the strong $CP$ problem is assumed. For definiteness, we restricted the analysis to examining the paradigm minimal supergravity (mSUGRA or CMSSM) model\\cite{msugra}. We were guided in our analysis by considering the possibility of including a viable mechanism for baryogenesis in the early universe. In order to do so, we needed to allow for re-heat temperatures after the inflationary epoch to reach values $T_R\\agt 10^6$ GeV. We found that in order to sustain such high re-heat temperatures, as well as generating predominantly {\\it cold} dark matter, we were pushed into mSUGRA parameter space regions that are very different from those allowed by the case of thermally produced neutralino dark matter. In addition, we found that very high values of the PQ breaking scale $f_a/N$ of order $10^{11}-10^{12}$ GeV were needed, leading to the mSUGRA model with {\\it mainly axion cold dark matter}, but also with a small admixture of thermally produced axinos, and an even smaller component of warm axino dark matter arising from neutralino decays. The favored axino mass value is of order 100 keV. We note here recent work on models with dominant axion CDM explore the possibility that axions form a cosmic Bose-Einstein condensate, which can allow for the solution of several problems associated with large scale structure and the cosmic background radiation\\cite{pierre}.  In this paper, we will examine the mSUGRA model under the assumption 1. of neutralino CDM and 2. that mixed axion/axino DM ($a\\tilde{a}$DM) saturates the WMAP measured abundance\\footnote{The possibility of mixed $a\\tilde{a}$CDM was suggested in the context of Yukawa-unified SUSY in Ref. \\cite{mix}.}. To compare the two DM scenarios, we will evaluate a measure of fine-tuning in the relic abundance \\be \\Delta_{a_i} \\equiv\\frac{\\partial\\log\\Omega_{DM}h^2}{\\partial\\log a_i} \\ee with respect to variations in fundamental parameters $a_i$ of the model. Such a measure of relic abundance fine-tuning was previously calculated in Ref.~\\cite{eo} in the context of just neutralino dark matter. Here, we will expand upon this and also consider fine-tuning of the relic density in the case of mixed $a\\tilde{a}$DM. Our main conclusion is that the relic abundance of DM is {\\it much less fine-tuned in the case of mixed $a\\tilde{a}$CDM, as compared to neutralino CDM}. Thus, we find that mixed $a\\tilde{a}$CDM is {\\it theoretically preferable to neutralino CDM}, at least in the case of the mSUGRA model, and probably also in many cases of SUGRA models with non-universal soft SUSY breaking terms.  We will restrict our work to cases where the lightest neutralino $\\tz_1$ is either the LSP or the next-to-lightest SUSY particle (NLSP) with an axino LSP; the case with a stau NLSP and an axino LSP has recently been examined in Ref. \\cite{fstw}. Related previous work on axino DM in mSUGRA can be found in Ref. \\cite{cmssm}.  The remainder of this paper is organized as follows. In Sec. \\ref{sec:inoDM}, we calculate the neutralino relic abundance fine-tuning parameter $\\Delta_{\\tz_1}$ in the mSUGRA model due to variation in parameters $m_0$ and $m_{1/2}$. We find, in good agreement with Ref. \\cite{eo}, that the WMAP allowed regions are all finely-tuned for low values of $\\tan\\beta$. For much higher $\\tan\\beta\\sim 50$, the fine-tuning is much less with respect to $m_0$ and $m_{1/2}$, but nevertheless high with respect to $\\tan\\beta$. In Sec. \\ref{sec:inoprob}, we review the gravitino problem, leptogenesis and the cosmological production of axion and axino dark matter. In Sec. \\ref{sec:axDM}, we calculate the fine-tuning parameter $\\Delta_{a\\ta}$ for mixed $a\\tilde{a}$CDM under the assumption of a very light axino with $m_{\\ta}\\sim 0.1-1$ MeV. The fine-tuning is always quite low, for both cases of mixed axino/axion CDM and mainly axion CDM. In the case of mainly axino CDM, we find the scenario less well-motivated since for high values of $T_R\\agt 10^6$ GeV, the value of $m_{\\ta}\\ll 0.1$ MeV, making the axino mainly {\\it warm} DM instead of cold DM. In Sec. \\ref{sec:conclude}, we present a summary and conclusions. ",
        "conclusions": "\\label{sec:conclude}  In this paper, we have examined the fine-tuning associated with the relic density of dark matter in the minimal supergravity model. We have calculated a measure of fine-tuning assuming two scenarios for SUSY dark matter: 1. neutralino dark matter with fine-tuning parameter $\\Delta_{\\tz_1}$, and 2. mixed axion/axino dark matter with fine-tuning parameter $\\Delta_{a\\ta}$.  In the case of neutralino dark matter, we find that the WMAP-allowed regions of mSUGRA such as the stau co-annihilation region, the HB/FP region and the light Higgs $h$-resonance annihilation region, are all rather highly fine-tuned, especially the HB/FP region, where $\\Delta_{\\tz_1}$ ranges from 20-100. Only mild fine-tuning is found in the low $m_0$, low $m_{1/2}$ region where stau co-annihilation and bulk annihilation through $t$-channel slepton exchange overlap. If one moves to large $\\tan\\beta\\sim 50$, then larger regions of parameter space which are consistent with WMAP occur. These large $\\tan\\beta$ regions have modest fine-tuning versus $m_0$ and $m_{1/2}$, but very large fine-tuning versus $\\tan\\beta$.  If instead we assume that dark matter is composed of an axion/axino admixture, rather than neutralinos, then we find that the relic density fine-tuning parameter is generically much lower: $\\Delta_{a\\ta}\\sim 1.3-2.5$ throughout parameter space. Here, we have assumed the existence of a light axino with mass $m_{\\ta}\\sim$ keV-MeV. Such a light axino opens up all of mSUGRA parameter space to being WMAP allowed, since now the neutralino decays via $\\tz_1\\to \\ta\\gamma$. If the DM is dominated by thermally produced axinos, then the re-heat temperature $T_R$ is generally lower than $10^6$ GeV unless the axinos are actually warm dark matter ($m_{\\ta}\\alt 100$ keV), so this scenario seems rather unlikely. However, if the PQ breaking scale $f_a/N$ is large, then the DM can be either a nearly equal axion/axino admixture, in which case fine-tuning is lowest ($\\Delta_{a\\ta}\\sim 1.3$), or a dominantly axion mixture (in which case $\\Delta_{a\\ta}\\sim 2.3$). Either scenario easily admits $T_R>10^6$ GeV, which can allow for non-thermal leptogenesis to occur.  The consequences of the mixed $a\\tilde{a}$CDM scenario for future dark matter searches is as follows. For collider searches, we expect much the same collider signatures as in the mSUGRA model with neutralino dark matter, since we assume the $\\tz_1$ is the NLSP, and decays far outside the collider detectors. However, {\\it all} of mSUGRA parameter space is now WMAP-allowed, instead of just the special co-annihilation, HB/FP region and resonance annihilation regions. As shown in Ref. \\cite{axdm}, the regions of WMAP-allowed neutralino CDM yield the lowest values of $T_R$, and so the stau co-annihilation, HB/FP region and $h$ resonance annihilation regions are most dis-favored for the case of mixed $a\\ta$CDM.  As far as WIMP searches go, in the mixed $a\\tilde{a}$CDM scenario, we expect no positive signals if $m_{\\tz_1}>m_{\\ta}$. If $m_{\\ta}>m_{\\tz_1}$, then the $\\tz_1$ would still be stable (assuming $R$-parity conservation) and WIMP direct and indirect detection signals are still possible\\cite{njp}. In the case of large axion relic abundance, which appears to us to be the favored scenario, then a positive signal at relic axion search experiments such as ADMX might be expected\\cite{admx}, although solar axion searches are less likely to achieve positive results, since large values of $f_a/N$ are favored, leading to small axion/axino couplings.  Our analysis has been based on the admittedly subjective basis of fine-tuning of the relic density of dark matter relative to model input parameters. We note here that the mSUGRA model already needs substantial fine-tuning in the electroweak sector in order to accomodate the relatively light $Z$ boson mass in the face of limits on the soft SUSY breaking parameters\\cite{ewft} (the little hierarchy problem). Our philosophy here is that less fine-tuning is better, and high fine-tuning in one sector is better than high fine-tuning in two sectors, {\\it e.g.} electroweak and dark matter sectors.  While our analysis has been restricted to the mSUGRA SUSY model, one might ask how general our conclusions might be. In SUSY models based on gravity mediation, with a neutralino LSP, the DM relic density is {\\it generically too high} unless some special mechanism is acting to enhance the neutralino annihilation cross section in the early universe.\\footnote{ Discussion on numerous different SUGRA models with non-universality has been explored in Ref. \\cite{wtn}.} For instance, in SUSY models with non-universality, instead of stau or stop co-annihilation, one might have sbottom or sneutrino co-annihilation, or bino-wino co-annihilation: in any case, the mass gap between co-annihilating particles must be fine-tuned to obtain agreement with the measured dark matter abundance. In non-universal models with a {\\it well-tempered neutralino}\\cite{wtn}, where the neutralino bino-higgsino or bino-wino composition is adjusted to fit the measured relic density, other parameters (Higgs soft masses, gaugino masses) must be fine-tuned to get just the right ``tempering'', as occurs in the mSUGRA HB/FP region. In other models, Higgs soft mass terms can be adjusted to allow $2m_{\\tz_1}$ to sit atop the $A$ resonance; but again, in this case, parameters must be fine-tuned (unless $\\tan\\beta$ is large, which also occurs in mSUGRA). The case where the SUSY neutralino abundance is not fine-tuned has long been noted: it is where squarks and sleptons are so light that $t$-channel annihilation channels are large. However, LEP2 search limits now essentially exclude all these regions. Thus, although we restrict our analysis here to the mSUGRA model, we feel this model provides a sort of microcosm for general SUSY models, in that it illustrates many of the features common to all SUSY models.  Our main conclusion is this. In the world HEP community, a tremendous effort is underway to explore for WIMP cold dark matter, based partly on the view that SUSY models naturally give rise to the ``WIMP-miracle'', and an excellent WIMP candidate for CDM. We have shown here that at least for the paradigm SUSY model-- mSUGRA-- usually a large overabundance of neutralino CDM is produced, unless one lies along a region of very high fine-tuning, where a slight change in model parameters leads to a large change in relic density: this equates to a high degree of relic density fine-tuning. Alternatively, if one assumes the PQWW solution to the strong CP problem within SUSY models, and a very light axino with $m_{\\ta}$ of the order of MeV, then along with an elegant solution to the strong CP problem, one obtains a mixed axion/axino relic density with much less fine-tuning. Given our results, we would advocate that a much increased share of HEP resources be given to relic axion searches, where the global search effort has been much more limited."
    },
    "0910/0910.5720_arXiv.txt": {
        "abstract": "We present a new analysis of stellar mass functions in the COSMOS field to fainter limits than has been previously probed at $z \\leq 1$.  The increase in dynamic range reveals features in the shape of the stellar mass function that deviate from a single Schechter function.  Neither the total nor the red (passive) or blue (star-forming) galaxy stellar mass functions can be well fit with a single Schechter function once the mass completeness limit of the sample probes below $\\sim 3 \\times 10^{9} \\Msun$. We observe a dip or plateau at masses $\\sim 10^{10} \\Msun$, just below the traditional $M^*$, and an upturn towards a steep faint-end slope of $\\alpha \\sim -1.7$ at lower mass at all redshifts $\\leq 1$. This bimodal nature of the mass function is {\\em not} solely a result of the blue/red dichotomy.  Indeed, the blue mass function is by itself bimodal at $z \\sim 1$.  This suggests a new dichotomy in galaxy formation that predates the appearance of the red sequence. We propose two interpretations for this bimodal distribution.  If the gas fraction increases towards lower mass, galaxies with $M_{\\rm baryon} \\sim 10^{10} \\Msun$ would shift to lower stellar masses, creating the observed dip.  This would indicate a change in star formation efficiency, perhaps linked to supernovae feedback becoming much more efficient below $\\sim$10$^{10} \\Msun$. Therefore, we investigate whether the dip is present in the baryonic (stars+gas) mass function.  Alternatively, the dip could be created by an enhancement of the galaxy assembly rate at $\\sim$10$^{11} \\Msun$, a phenomenon that naturally arises if the baryon fraction peaks at $M_{\\rm halo} \\sim 10^{12} \\Msun$. In this scenario, galaxies occupying the bump around $\\Mstar$ would be identified with central galaxies and the second fainter component of the mass function having a steep faint-end slope with satellite galaxies. The low-mass end of the blue and total mass functions exhibit a steeper slope than has been detected in previous work that may increasingly approach the halo mass function value of -2. While the dip feature is apparent in the total mass function at all redshifts, it appears to shift from the blue to red population, likely as a result of transforming high-mass blue galaxies into red ones.  At the same time, we detect a drastic upturn in the number of low-mass red galaxies.  Their increase with time seems to reflect a decrease in the number of blue systems and so we tentatively associate them with satellite dwarf (spheroidal) galaxies that have undergone quenching due to environmental processes. ",
        "introduction": "\\label{sec:introduction}  Galaxy formation and evolution is believed to be driven primarily by two processes: firstly, the successive merging of their parent dark matter halos causing accretion of material and ultimately mergers between galaxies; and secondly, the feedback-regulated conversion of gas into stars within galactic disks with subsequent potential rearrangement of the disk material by dynamical processes (secular evolution).  Both processes contribute to the growth in stellar mass of galaxies with time. The stellar mass function of galaxies and its evolution with time is therefore fundamental to the understanding of galaxy formation.  The ability to estimate galaxy stellar masses has advanced in recent years in large part because of increasing access to near-IR photometry.  Estimates are typically made by fitting multi-band photometry with stellar population synthesis libraries (see, e.g., \\citealp{Brinchmann+2000,Bruzual+2003,Drory+2003,Drory+2004,Maraston+2006,Marchesini+2008,Conroy+2009}), fitting specific spectral features when spectroscopy is available \\citep{Kauffmann+2003}, or the full spectrum when high-quality spectra are observed \\citep{Reichardt+2001,Panter+2004}. So far, only the photometric fitting technique has been a viable option for high-redshift surveys. These measurements provide masses with accuracies of $\\sim 0.1-0.3$~dex, and systematic uncertainties of up to a factor of two depending on the selection and number of photometric bands included and the assumptions made on, among others, the shape of the IMF, the allowed star formation histories, the dust extinction model, or the underlying stellar population synthesis method (see, e.g., \\citealp{Drory+2004,Kannappan+2007,Marchesini+2008,Conroy+2009} for systematic studies on this matter).  Utilizing such techniques, the build-up of the stellar mass density from redshift $z \\sim 6$ to the present epoch has been the subject of several studies in the past decade, often relying on deep multi-band imaging surveys in the UV to near-infrared wavelength range \\citep{Brinchmann+2000,Drory+2001a,Cohen2002,Dickinson+2003,Fontana+2003,Rudnick+2003,Glazebrook+2004,Conselice+2005,Chapman+2005,Drory+2005a,Rudnick+2006,Eyles+2007,Grazian+2007,Stark+2007}.  The stellar mass function in the local universe has been measured from large imaging and spectroscopic surveys such as 2dF, SDSS, and 2MASS \\citep{Cole+2001,Bell+2003,Perez-Gonzalez+2003,Panter+2004,Baldry+2008}. At distances up to $z \\sim 1.5$, a number of groups have established a picture of the evolution of the mass function with some detail \\citep{Drory+2003,Fontana+2004,Bundy+2005,Borch+2006,Bundy+2006,Arnouts+2007,Pozzetti+2007,Ilbert+2009}, with generally good agreement between different datasets. To some lesser detail and accuracy, deep surveys have provided data spanning $0 < z \\lesssim 5$ \\citep{Drory+2005a,Conselice+2005,Fontana+2006,Yan+2006,Grazian+2007,Elsner+2008,Perez-Gonzalez+2008,Marchesini+2008}, and even some estimates at $z \\sim 7$ \\citep{Bouwens+2006}. So far, these high-redshift studies of the stellar mass function emphasized the evolution of galaxies of mass $\\logMsun \\gtrsim 10$. Speaking very broadly, the stellar mass density decreases by a factor of two to $z \\sim 1$, with the most massive galaxies already being in place at earlier epochs. The evolution appears to accelerate quickly beyond $z \\sim 1.5$.  In this paper, we concentrate on the low-mass galaxies that have typically been below the completeness limits of previous work at $z > 0.1$. Generally \\citet{Schechter1976} fits to the galaxy stellar mass function with faint-end slope $\\sim -1.1$ to $\\sim -1.3$ have been found adequate to describe the galaxy population (even separated morphologically, by color, or star formation activity; \\citealp{Pannella+2006,Borch+2006,Arnouts+2007,Pannella+2009,Ilbert+2009}). Recently, though, a steepening of the slope of the luminosity function below $M_i \\sim -17$ in the local universe has been convincingly detected in clusters \\citep{Driver+1994a,Trentham+2002,Hilker+2003,Popesso+2005,Popesso+2006}, groups \\citep{Trentham+2002,Trentham+2005,Gonzalez+2006}, and in the field \\citep{Blanton+2005a}. For example, \\citet{Trentham+2002} find that the luminosity function in the Virgo cluster, in the NGC~1407 group, and in the Coma~1 group is steep between $M_R$ of -18 and -15 (and flattens again only at $M_R > -15$). Moreover, \\citet{Baldry+2008} find that the local galaxy stellar mass function steepens as well below $\\logMsun \\sim 9.5$ (but see also \\citealp{Li+2009}).  The steepening of the mass function can also be interpreted as a bimodality: the mass function consists of a sum of (at least) two components. This bimodal behavior has now also been detected at redshifts $z>0.1$. \\citet{Pozzetti+2009} find bimodal mass functions to $z\\sim 0.5$ from the zCOSMOS spectroscopic survey. They interpret the mass function as being composed of early-type galaxies dominating the massive part and late-type galaxies dominating the less massive part and contributing the steep faint-end slope. Each of these components is well fit by a Schechter function. \\citet{Bolzonella+2009} use the same sample to investigate the bimodality as a function of environment. They find that at $z\\lesssim 0.5$, the shape of the galaxy stellar mass function in high- and low-density environments become markedly different, with high density regions showing a stronger bimodality.  We extend the study of the shape of the galaxy stellar mass function, particularly at low masses, to $z \\sim 1$, with stellar mass limits $\\sim 1.5$~dex lower than can be achieved with spectroscopic studies. We investigate actively star-forming galaxies and passively evolving galaxies separately; we study how the change in slope may be caused by the presence of multiple galaxy populations, that taken together, lead to a mass function shape that is more complex than a single power law with an exponential cutoff or even a simple combination of early- and late-type components. We show that the blue mass function itself is bimodal and that passive galaxies exhibit a faint-end upturn, likely caused by dwarf spheroidal galaxies linked to the faint end of the blue galaxy population.  This paper is organized as follows: in \\S~\\ref{sec:sample} we introduce the galaxy sample that we use in this work. In \\S~\\ref{sec:stellar-mass} we discuss the stellar population models used to derive stellar masses and the resulting mass completeness limits. In \\S~\\ref{sec:mf} we present the stellar mass function of active and passive galaxies and we discuss our results in \\S~\\ref{sec:discuss}. Finally, we summarize this work in \\S~\\ref{sec:summary}.  Throughout this work we assume $\\Omega_M = 0.3$, $\\Omega_{\\Lambda} = 0.7, H_0 = 70\\ h_{70}^{-1} \\mathrm{km\\ s^{-1}\\ Mpc^{-1}}$. Magnitudes are in the AB system. We will denote galaxy stellar masses by the symbol $M$ -- or $\\Mg$ where an explicit distinction form halo masses, denoted by $\\Mh$, is necessary. The symbol $M^*$ is reserved for the characteristic mass parameter of the Schechter function. ",
        "conclusions": "\\label{sec:discuss}   \\subsection{Comparison to Previous Results}  \\begin{figure*}[ht] \\centering \\includegraphics[width=0.8\\textwidth]{fig8.eps} \\caption{A comparison of the mass function deduced in this work with data from the literature. The mass function of passive galaxies is marked in red, that of star-forming galaxies in blue, and that of all galaxies in black. The solid lines show double Schechter function fits to the data (Eq.~\\ref{eq:schechter2}; the thin lines show the individual bright and faint components of the double Schechter function). The fit at low $z$ is repeated in higher-$z$ panels as dotted lines. The magenta line denotes the \\citet{Baldry+2008} $z \\sim 0$ SDSS-based mass function convolved with the photo-z error distributions (see text). The dashed blue, red, and black lines are the 3.6$\\mu$m-selected mass functions of active, passive, and all galaxies, respectively, in the COSMOS survey by \\citet{Ilbert+2009}. The green dash-dotted line is the mass function by \\citet{Perez-Gonzalez+2008}. Additionally, we plot the Press-Schechter halo mass function scaled by the global baryon fraction $f_{\\rm b}$ (\\citealp{Dunkley+2009}; WMAP5) in cyan. \\label{fig:mf_lit}} \\end{figure*}  Before discussing possible interpretations of the more complicated mass function shape apparent in the COSMOS data, it is useful to compare our results to previous work with the aim of determining whether evidence for similar behavior has been found in other surveys. In Fig.~\\ref{fig:mf_lit} we compare our mass function to a number of results from the literature and for later reference we include the halo mass function \\citep{Reed+2007} where the halo masses have been multiplied by the global baryon fraction, $f_{\\rm b} = \\Omega_{\\rm b}/\\Omega_{\\rm m}$, taken from the WMAP 5-yr data \\citep{Dunkley+2009}. We show the $z \\sim 0$ SDSS-based mass function by \\citep{Baldry+2008} convolved with the photo-z error distributions as shown in Fig.~\\ref{fig:mc_conv} and discussed in \\S~\\ref{sec:photoz_uncertainties}. We also plot the mass functions in the redshift range $0<z<1$ of active, passive, and all galaxies in the COSMOS survey selected at 3.6$\\mu$m by \\citet{Ilbert+2009} using the same photometric redshifts as the present work.  Finally the total mass function selected at $3.6-4.5\\mu$m by \\citet{Perez-Gonzalez+2008} is also plotted.  Where they overlap in mass, our mass functions agree very well with previous work, with some systematic differences stemming from differences in the underlying stellar population synthesis libraries and grids of star formation histories (see \\citealp{Marchesini+2008} and \\citealp{Conroy+2009} for a systematic study on the influence of stellar population grid parameters on mass determinations). Of notable importance, though, is the fact that the bright components in our two-component fit to the mass function agree very well with the blue and the red sub-components in \\citet{Ilbert+2009}, despite (minor) differences in the definition of blue and red galaxies. This gives us confidence that the decomposition of the mass function into a bright and a faint component with six free parameters is not strongly degenerate. Specifically, both the mass scale of the bright components and their faint-end extrapolations are compatible with the single Schechter function fits to the more restricted data sets. We also confirm the build-up of the faint part of the red sequence at $z < 1$ observed by other groups \\citep[e.g.][]{Bell+2004,Bundy+2006,Bell+2007,Faber+2007,Perez-Gonzalez+2008,Ilbert+2009,Williams+2009}.  As we discuss further below, the apparent multi-component nature of the mass function has been discussed at low redshifts by \\citet{Baldry+2008} and \\citet{Li+2009} and observed at some level in other high-$z$ spectroscopic and hence significantly shallower studies of the COSMOS field \\citep{Pozzetti+2009, Bolzonella+2009}.  The advantage of the current analysis is that it probes to lower mass at redshifts beyond $z = 0.2$, thus allowing a larger dynamic range which for the first time at high-$z$ reveals more complicated behavior, specifically the steepening of the faint-end slope and an additional population of faint red galaxies (see also \\citealp{Salimbeni+2008}). A similar steepening in the luminosity function below $M_i \\sim -17$ has been convincingly detected recently in clusters \\citep{Driver+1994a,Trentham+2002,Hilker+2003,Popesso+2005,Popesso+2006}, groups \\citep{Trentham+2002,Trentham+2005,Gonzalez+2006}, and in the field \\citep{Blanton+2005a}. \\citet{Baldry+2008} find that the local galaxy stellar mass function steepens as well below $\\logMsun \\sim 9.5$ (see also \\citealp{Salucci+1999}).   \\subsection{Halo Mass to Stellar Mass Relation}  \\begin{figure}[hbt] \\vspace*{0.4cm} \\centering \\includegraphics[width=8cm]{fig9.eps} \\caption{Stellar mass as a function of halo mass (bottom panel) and ratio between stellar mass and halo mass (top panel) at redshifts of $z=0.3$ and $z=0.9$ determined by abundance matching the galaxy stellar mass function to the halo mass function. The relation determined by \\citet{Moster+2009} are shown for comparison. \\label{fig:abmatch} } \\end{figure}  For the following discussion of the physical significance of the various morphological features of the mass function, it is useful to first place the mass function in the context of the halo mass function. We will do so by analyzing the halo mass vs.\\ galaxy stellar mass relation, noting that the more complicated shape of the stellar mass function also leads to a more complicated relation between galaxy stellar mass, \\Mg, and halo mass, \\Mh.  The relationship between the halo mass function and the galaxy (stellar) mass function can be written as \\begin{equation} \\phig(\\Mg,t)=\\left|\\frac{d\\Mh}{d\\Mg}\\right|\\phih(\\Mh,t). \\label{eq:abmatch0} \\end{equation} Such a one-to-one correspondence between halo mass and galaxy stellar mass at some time, $\\Mh(\\Mg)$, can be found by requiring that the cumulative number density of halos above a given mass and galaxies above a corresponding stellar mass be equal (abundance matching), \\begin{equation} \\int_{\\Mg}^\\infty \\phig(M)dM = \\int_{\\Mh}^\\infty \\phih(M)dM. \\label{eq:abmatch} \\end{equation} This approach has been shown to be sufficiently accurate to match the two-point correlation function of galaxies and halos \\citep{Conroy+2006,Moster+2009}.  To obtain a more realistic picture of the halo masses of galaxies, we must include sub-halos in our abundance matching procedure, as halos (and hence their galaxies) may survive for considerable time. Analyzing high-resolution cosmological N-body simulations, \\citet{Giocoli+2008} and \\citet{Angulo+2008} provide fitting formulae for the number of sub-halos per host halo per logarithmic interval of sub-halo to host mass ratio, $d\\Ns/d\\ln(\\Ms/M)$ (see also \\citealp{Gao+2004}). This quantity can be converted to a sub-halo mass function, $\\phi_{\\rm sub}(\\Ms)$, which is the more convenient quantity for our purposes by changing variables to obtain $d\\Ns/d\\Ms$, multiplying by the abundance of host halos, and integrating over host halo mass from $\\Ms$ to infinity: \\begin{equation} \\phi_{\\mathrm sub}(\\Ms) = \\frac{d\\Ns}{d\\Ms} = \\int_{\\Ms}^\\infty \\frac{dN_h}{d\\Mh} \\frac{1}{\\Ms} \\frac{d\\Ns}{d\\ln(\\Ms/M)} dM. \\label{eq:subhalomf} \\end{equation} This sub-halo mass function can then be added to the distinct halo mass function to obtain the total mass function that we use to match to the abundance of galaxies: \\begin{equation} \\phi(M) = \\phi_{\\rm distinct} + \\phi_{\\rm sub}. \\label{eq:totalhalomf} \\end{equation}  We note that we use the mass of the sub-halo at infall time, not the mass of the sub-halo at the time of observation for matching abundances. Most of the galaxy's properties will have been set by the time infall occurs (and shaped by the potential of the halo at that time) and galaxy growth is likely to slow down or even stop as soon as the galaxy becomes a satellite. This approach is motivated by models of galaxy formation that generically predict that stellar mass is tightly linked to the potential well in which the galaxy forms \\citep{Kauffmann+1993,Cole+1994,Somerville+1999}. Sub-halos loose mass by tidal stripping, however their stars are more centrally concentrated and are not stripped until most of the dark matter has been lost. The relevant mass scale for sub-halos, therefore, is the virial mass at infall time.  We proceed using a halo mass function consisting of the distinct halo mass function by \\citet{Reed+2007} and the sub-halo abundance by \\citet{Giocoli+2008} processed through Eq.~\\ref{eq:subhalomf}. In Fig.~\\ref{fig:abmatch}, we show the resulting relation between halo mass, \\Mh, and galaxy stellar mass, \\Mg at redshifts $z=0.3$ and $z=0.9$ obtained by abundance matching to our total galaxy stellar mass function. For comparison, we also show the relation obtained by \\citet{Moster+2009} by populating halos from N-body simulations with galaxies using semi-analytic methods and the requirement that the stellar mass function be reproduced.  Without accounting for substructure, $\\Mg(\\Mh)$ evolves relatively uniformly at all masses.  With substructure, much less change in the relation at low mass is seen; there is little evolution in $\\Mg(\\Mh)$ from $z=0.3$ to $z=0.9$ at $\\log \\Mh/\\Msun \\lesssim 11.5$ (corresponding to $\\log \\Mg/\\Msun \\lesssim 9.5$, about the mass of the LMC). This is to be expected if the stellar content of low-mass halos is limited by feedback. There is evolution in the $\\Mg(\\Mh)$-relation at masses above $\\Mstar$, where the stellar mass per halo mass increases by a factor of 1.23 from $z=0.9$ to $z=0.3$. The peak in star formation efficiency occurs in $\\logMsun \\sim 12.2$ halos and reaches 23\\% at $z=0.3$ and 19\\% at $z=0.9$; galaxy formation -- if there is a monotonic relation between stellar mass and halo mass such as the one arrived at by abundance matching -- is always an inefficient process.  While the flattening of the $\\Mg(\\Mh)$-relation above $\\Mstar$ (corresponding to the exponential cut-off in the galaxy mass function) is understood as an effect of the inefficiency of cooling in large halos (e.g.\\ \\citealp{Rees+1977,White+1978}), the steepening of the $\\Mg(\\Mh)$-relation at halo mass of $\\logMsun \\sim 11$ suggests another qualitative change in behavior of the feedback efficiency.  The $\\Mg(\\Mh)$-relation is fundamental to the process of galaxy formation within dark matter halos and understanding what shapes this relation is equivalent to understanding the galaxy stellar mass function. At this point we wish to remark that, due to the derivative factor $d\\Mh/d\\Mg$ in Eq.~\\ref{eq:abmatch0}, or equivalently, the integral nature of the abundance matching relation in Eq.~\\ref{eq:abmatch}, a single change of slope of $\\Mg(\\Mh)$ leads to a dip-bump structure in the inferred galaxy mass function, given a power-law halo mass function.  The more abruptly the change of slope occurs, and the larger it is, the more pronounced the dip-bump feature in the mass function becomes. The mass function shape we observe, therefore, leads to an approximately double-power-law form of the $\\Mg(\\Mh)$ relation, as can be seen in Fig.~\\ref{fig:abmatch}.   \\subsection{Interpreting the Shape of the Mass Function}\\label{sec:interpretation}  There are several ways to interpret the shape of the total and type-dependent mass functions in COSMOS.  Broadly speaking, we will concentrate on two interesting features, namely the ``upturn'' towards steeper faint-end slopes at low masses ($\\logMsun \\sim 9$) and an apparent ``dip'' or plateau at masses just above $\\logMsun \\sim 10$ that sets in below the traditional $M^*$.  We note that even breaking down the mass function into these two components is, in itself, an interpretation since the steepening faint-slope alone may provide a viable explanation.  As we discuss below, our interest in separately identifying the dip is motivated by comparisons to the underlying halo mass function which is a strict power-law at the relevant masses.  In this context, the dip represents a suppression of the stellar mass associated with halos in this range; this is a consequence of a change in $d\\Mh/d\\Mg$ such that there is a smaller change in the mass of galaxies per change in the mass of halos. The purpose of this section is, hence, to explore several physical mechanisms that may explain both the upturn and the dip and to speculate on the role of future observations in distinguishing them.  Using the zCOSMOS spectroscopic survey \\citet{Pozzetti+2009} interpret the mass function they recover as being composed of early-type galaxies dominating the massive part and late-type galaxies dominating the less massive part and contributing to the steep faint-end slope. They find that each of these components is well fit by a Schechter function in the mass range they consider. \\citet{Bolzonella+2009} use the same sample to investigate the bimodality as a function of environment. The find that at $z\\lesssim 0.5$, the shape of the galaxy stellar mass function in high- and low-density environments become extremely different, with high density regions showing a stronger bimodality. Also, \\citet{Ilbert+2009} remark that the 3.6$\\mu$m-selected COSMOS mass function of all galaxies is not well-fit by a single Schechter function. They prefer the sum of their (Schechter) fits to the red and blue galaxies to describe the whole population, as this reproduces the dip at intermediate mass of $\\log M \\sim 10$ and $z \\lesssim 0.6$ they observe. Due to the IRAC selection, they do not reach faint enough limits to detect the steep faint-end slope or the bimodal structure at higher $z$. We over plot the data from the FDF \\citep{Drory+2005a} at $z=0.5$ in Fig.~\\ref{fig:mf_data} and note that the bimodality is visible there as well, although \\citet{Drory+2005a} do not discuss it because a single Schechter fit seemed adequate at the time (apart from the data point at $\\log M = 11$ which comes out too high).  We emphasize that we see even more structure in the mass function than observed in the works discussed above; those studies only detect a dip in the total mass function and interpret its origin as being due to the superposition of active and passive populations which by themselves are unimodal. In this paper, we detect a significant change in slope in the mass functions of active galaxies as well as a marked bimodality in the mass function of passive galaxies.  As mentioned above, previous work has attributed the two-component nature of the mass function to a superposition of a (unimodal) blue galaxy population modeled as a Schechter function and a red galaxy population, also modeled as a Schechter function but with a larger characteristic mass.  In this picture, the bimodality arises through the transformation of blue galaxies into red galaxies in a process that must be linked to the mass evolution of a transitioning galaxy to deplete the dip and create a bump at $M \\sim M^*$.  In contrast, we find that the blue galaxy population by itself shows a dip signature and a steepening of the power-law slope, which can be interpreted as arising from a bimodal distribution. In fact, both the red and the blue populations in Fig.~\\ref{fig:mf_data} can be interpreted as bimodal. However, the bimodality in the blue population seems to be more pronounced at high $z$ compared to low $z$: it becomes weaker as redshift decreases. This effect can also be seen in the $\\chi^2$ values reported in Table~\\ref{tab:mffit}. Indeed, the best $\\chi^2$ values for single Schechter fits are obtained for the blue component at low redshift ($\\chi^2_{\\rm single} = 7.0$). The signature of the bimodality (the dip at intermediate mass, or the bump at around $M^*$) ``moves'' from the blue population at high $z$ to the red population at low $z$.  This is likely because some of the blue galaxies that make up the bump around $M^*$ in the blue mass function turn red with time \\citep[e.g.][]{Bell+2004,Bundy+2006,Faber+2007,Williams+2009}.  The fact that the blue mass function is itself bimodal at $z \\sim 1$ implies that the shape of the total mass function predates the emergence of the red sequence.  We therefore suggest that the physical explanation must be a mass-dependent effect or mechanism that is largely separate from the process that transforms star-forming disks into passive spheroidals.  In other words, this suggests a new dichotomy in the galaxy distribution besides the well-known red--blue distinction.  We will follow this hypothesis in the following, using the behavior of blue galaxies to stand in for interpretations of the full population and returning to red systems in the subsequent section.     \\subsubsection{Blue Galaxies}\\label{sec:blue}  Before discussing possible physical explanations, we note that among star-forming, blue galaxies, there is an interesting separation between (giant) spirals of Hubble types Sa--Scd that dominate the mass function at the massive end and make up the majority of objects above the dip, and dwarf galaxies of Hubble types Sd--Sm/Im that dominate the less massive, power-law part of the blue galaxy mass function. Disk parameters such as disk radius and luminosity remain roughly constant (with a large spread) along the Hubble sequence at types Sa-Scd. For types later than Sd galaxies, the Hubble sequence turns into a luminosity sequence, with later type morphologies corresponding to smaller and less massive galaxies (see, for example, the review by \\citealp{Roberts+1994} and references therein).  It is also worth noting that Disk galaxies without a significant bulge component exist (e.g., Hubble type Sd and later) and can be accurately characterized as pure disk systems, harboring at most a nuclear star cluster.  May we tentatively identify this distinction in disk properties with the two components in the blue population's mass function?  If so, the physical processes discussed below must also account for these morphological differences.  {\\bf Star Formation Efficiency.}  The dip in the stellar mass function corresponds to a change in the stellar mass vs.\\ halo mass relation such that there is a smaller change in the number of galaxies when measured against the power-law shape of the halo mass abundance (Figs.~\\ref{fig:mf_lit} and \\ref{fig:abmatch}).  It is possible that this deficit arises from a change in the efficiency of star formation with mass.  In this interpretation, the {\\em baryonic} mass function, if it could be observed, would show a close to power-law shape below $M^*$, indicating a more smoothly varying fraction of $M_{\\rm baryon} / M_{\\rm halo}$ with $M_{\\rm halo}$. For this to still be the case at later cosmic times we must also require that the fraction of gas retained by any halo in the interesting mass range must be a smooth function of halo mass and also that this fraction not be too small. Then, a decrease in SFR efficiency below a given mass scale (say $\\logMsun \\sim 10$) would lead to an increasing gas fraction, $f_{\\rm gas}$, at lower masses and a {\\em decreasing} stellar mass fraction. This is analogous to the break in the Tully-Fisher relation \\citep{Tully+1977} where galaxies with $V_{\\rm c} \\lesssim 90$~km~s$^{-1}$ fall below the relation defined by more massive galaxies. If plotted against total disk mass, $M_{\\rm gas}+M_{\\rm star}$, instead of luminosity, a single relation is restored with $M_{\\rm gas}+M_{\\rm star} \\sim V_{\\rm c}^4$ \\citep{McGaugh+2000}.  The decrease in $f_{\\rm gas}$ as a function of $M$ has been observed. Thus galaxies with baryonic masses near $\\logMsun \\sim 10$ would have observed $M$ several factors lower, thereby creating a dip in the stellar mass function at $\\logMsun \\sim 10$, and leading to a steepening below $\\logMsun \\sim 9$.  \\begin{figure} \\centering \\includegraphics[width=8cm]{fig10.eps} \\caption{The top panel shows the gas fraction, $f_{\\rm gas}$, as a function of stellar mass.  The data are taken from a compilation by \\citet{Hopkins+2009}, including \\citet[][ triangles]{Bell+2001},\\citet[][circles]{Kannappan2004}, and \\citet[][squares]{McGaugh2005}. A spline interpolation to the data is shown as a dashed line. The gas fraction necessary to remove the dip structure in the mass function is shown as a dash-dotted line.  In the bottom panel, we plot the total stellar mass function at $z=0.3$ and the baryonic mass function obtained by applying the gas fractions from the top panel. Power-law slopes of $\\phi \\sim M^{-2}$ for the halo mass function and $\\phi \\sim M^{-1.7}$ for the stellar mass function are plotted for reference. \\label{fig:gas_fraction}} \\end{figure}  Using a simple fitting spline to observations of $f_{\\rm gas}(M)$ compiled by \\citet{Hopkins+2009} including data from \\citet{Bell+2001}, \\citet{Kannappan2004}, and \\citet{McGaugh2005}, we transform the total stellar mass function at $z=0.3$ into an approximate baryonic mass function. The gas fraction hereby is defined as the ratio of gas mass to stellar plus gas mass, $f_{\\rm gas} = M_{\\rm gas}/(M_{\\rm gas}+M_{\\rm star})$. We plot the results in Fig.~\\ref{fig:gas_fraction}.  Using the observed gas fraction as a function of mass, the dip in the mass function can be lessened, but not removed. However, as expected, an appropriate {\\em assumption} for $f_{\\rm gas}(M)$ leads to baryonic mass functions in which the dip can be ``straightened'' out (dot-dash line).  However, the amount of $f_{\\rm gas}(M)$ required is about a factor of 2 higher than current low-$z$ observations (dashed line).  This comparison involves several large uncertainties, but plausibly, the difference could be made up for in a warm phase that is not easily detected \\citep{Cen+1999,Dave+2001,Cen+2006}.  One would also expect $f_{\\rm gas}$ to increase with redshift in line with rising star formation rates.  We are able to conclude that it is at least {\\em possible} that the dip in the stellar mass function reflects a mass dependence in the ability for baryons to cool and form stars.  What causes this mass-dependent SFR efficiency?  We note that the deviation of the stellar mass function from power-law form at $\\logMsun \\sim 9-10$ is close to the scale where supernova feedback becomes less important, $V_c \\sim 100$~km~s$^{-1}$ \\citep{Dekel+1986,Benson+2003}. As long as supernova feedback is efficient in regulating the conversion of gas into stars, the ratio of stellar mass to halo mass is set by the scaling relation of feedback efficiency with halo mass (or circular velocity). Once the (DM) potential is deep enough that supernova feedback can no longer remove significant amounts of gas from the halo, the attainable ratio of stellar mass to dark matter mass becomes larger as a larger fraction of the gas in the halo can cool to form stars. As a consequence, the stellar mass function deviates upwards from its power-law form. (At the massive end, the stellar mass is limited by the cooling timescale becoming too large and the ratio of stellar mass to dark mass decreases again sharply causing the exponential cut-off in the stellar mass function \\citep{Rees+1977,White+1978}.  The transition from SFR efficiency suppressed by supernovae feedback to peak SFR efficiency may therefore lead to the observed dip in the stellar mass function.  There are two potential tests for this scenario.  The most obvious is to seek better measures of $f_{\\rm gas}$ as a function of $M$ and redshift, but such observations remain challenging.  A second test is to determine whether the mass scale that defines the dip evolves with redshift.  Because we expect $f_{\\rm gas}$ to be higher in more massive galaxies at early times (since at $z > 1$ many more are likely to be rapidly forming stars), the dip in the stellar mass function should move to higher masses at early times as well.  Unfortunately, the mass functions presented here are likely to be too affected by cosmic variance to robustly measure redshift evolution in the dip feature, but future data sets will be able to overcome this problem. Also, a more rigorous analysis taking into account cosmological gas accretion as well as a realistic feedback prescription and outflow model is needed, but is beyond the scope of this paper.  {\\bf Hierarchical assembly.} Another possibility is that the dip at $\\log M \\sim 10$ and the bump around $M^*$ are formed by a depletion of stellar and gas mass at intermediate mass resulting from a pile-up of the end-products of mass-dependent assembly around $M^*$.  In this scenario, both the stellar and baryonic mass functions would exhibit dips at $\\log M \\sim 10$ because the assembly of baryonic mass is accelerated at a certain scale.  It is known that the halo-halo merger rate depends very weakly on mass \\citep{Fakhouri+2008}.  However, if the stellar mass, or in this case, baryonic mass fraction, peaks in halos at a given mass scale---perhaps as a result of the SFR feedback discussed above---the resulting merger rate in terms of stellar mass, $M$, or baryonic mass, $M_{\\rm baryons}$, can be enhanced near this scale \\citep{Bundy+2009,Stewart2009,Hopkins+2009a} and strengthen the dip feature.  This is because the dynamical friction timescale (and hence the merger timescale) depends firstly on the mass ratio of the merging systems (e.g.\\ \\citealp{Boylan-Kolchin+2008} and references therein). Low mass ratio mergers happen only after a long time, and one-to-one mergers happen quickly. At low mass, one-to-one mergers in DM correspond to minor mergers in baryons and hence the baryonic content of a halo assembles more slowly than the halo.  This is due to the steep relation between stellar mass, \\Mg, and halo mass, \\Mh, below $M^*$ (see Fig.~\\ref{fig:abmatch}).  At masses near and just above the scale at which $\\Mg(\\Mh)$ flattens, more frequent minor halo mergers can host one-to-one baryonic mergers and thus the baryonic assembly rates increase markedly.  At even higher mass scales, above the flattening at $M^*$ of the $\\Mh(\\Mg)$-relation, the infalling galaxy is much more likely to become and stay a satellite, even if the baryonic mass ratio is close to one.  At this point, the baryonic merger rate drops again.  As a result, low stellar (baryonic) mass galaxies get depleted more rapidly in one-to-one mergers, and pile-up around the $M^*$ mass scale ($M^* \\sim 10^11~\\Msun$). This scenario also suggests that the two populations of galaxies might be thought of as central galaxies occupying the bump with an increasing fraction of satellite galaxies at lower masses that possibly form the second, faint, component of the mass function, leading to a steepening of the low-mass end of the mass function.  Such satellites in a galaxy or group halo would orbit for a long time before merging and hence remain visible as individual galaxies.  If this enhanced assembly scenario were important, it would also indicate that many major mergers between star-forming (disk) galaxies must not lead to a final destruction of the disk and truncation of star formation as the bump signature is already apparent in the blue galaxy population at $z\\sim 1$. Such mergers must produce a remnant that is still a blue star-forming disk galaxy. Such outcomes of mergers seem possible if the gas fraction is large \\citep{Springel+2005a,Robertson+2006,Governato+2007,Hopkins+2009b}. Also, we would expect the dip feature to deepen with time but, in contrast with SFR scenario above, remain relatively fixed with respect to mass.  {\\bf Steepening Faint-End Slope.} So far we have focused on the dip in the stellar mass function.  We now turn to the steep slopes at the low-mass end.  To the extent that this feature is unrelated to the dip forming process mentioned above, we can consider several possibilities. To begin with, many authors interpret the faint-end slope as rising to the maximum set by the halo mass function (with a value of -2).  Such behavior is expected if feedback processes that regulate the star formation efficiency become increasingly independent of $M_{\\rm halo}$ below some scale.  At low masses, for example, supernova feedback may so dominate over the ability of halos to retain gas, that the effective gas or stellar mass fraction becomes a constant, independent of halo mass.    \\subsubsection{Red Galaxies}\\label{sec:red}  We now turn to the behavior of the red galaxy mass function.  In the redshift bins $z=0.5$ and $z=0.3$, where the faint red component of the mass function is sampled well, the number density of all galaxies at $M \\lesssim 9.5$ does not change much (Fig.~\\ref{fig:mf_lit} and Table~\\ref{tab:mffit}). Yet, the number density of faint red galaxies increases by 0.45~dex, a combined effect of the normalization increasing from $\\phif(z=0.5) = 0.28$ to $\\phif(z=0.3) = 0.49$ and the characteristic mass of the low-mass component increasing from $\\Mstarf(z=0.5) = 9.41$ to $\\Mstarf(z=0.3) = 9.54$. This evolution in faint red galaxies is nearly perfectly mirrored by a decrease in number of faint blue galaxies. As we have noted above, the faint end slope of the active and the passive population is consistent with being equal.  The similarity of the steep faint-end slopes of the active and faint passive populations and their reciprocal change in number density is suggestive of the latter originating from the former by shutting off star formation. It is very tempting to identify these two faint galaxy populations sharing the same steep faint-end slope with Sm/Im/dIrr galaxies, and faint spheroidal galaxies, respectively. It is established that passively evolving dwarf galaxies cluster around massive galaxies \\citep{Zehavi+2005,Haines+2006,Haines+2007,Carlberg+2009}, and tidal interactions or ram pressure stripping may well lead to quenching and some subsequent phase mixing turning Sm/Im galaxies into faint spheroidal galaxies \\citep{Mayer+2001,Grebel+2003,Haines+2007}. Also, \\citet{McCracken+2008} finds that in the CFHTLS in the redshift bin $0.2 < z < 0.6$ the clustering amplitude for faint red galaxies is actually higher than that of bright red galaxies, which is to be expected in our interpretation.      Following on previous studies of the stellar mass function in the COSMOS field \\citep{Ilbert+2009, Pozzetti+2009, Bolzonella+2009}, we present a new analysis of this data set that provides mass functions to fainter limits than has been previously probed at $z \\lesssim 1$. The resulting increase in dynamic range allows us to characterize and study features in the shape of the stellar mass function that deviate from a single Schechter function.  We have tested whether these features could be introduced by a variety of systematic effects including both catastrophic photometric redshift errors as well as increasing photo-$z$ uncertainty at low masses, differences in the way stellar population models account for TP-AGB stars, and the ability of stellar mass codes to convert from luminosity to mass.  We conclude that our results are robust to these effects, although the data are still limited by cosmic variance.  Our key results follow:  \\begin{itemize}  \\item Neither the total nor the red (passive) or blue (star-forming) galaxy stellar mass functions can be well fit with a single Schechter function once the the mass completeness limit of the sample probes below roughly $3 \\times 10^{9} \\Msun$.  We model this more complicated behavior using a double Schechter function with 6 free parameters, and present the fitting results for 4 redshift bins to $z=1$.  \\item The bimodal nature of the mass function is {\\em not} solely a result of the blue/red dichotomy.  Indeed, the blue mass function is already bimodal at $z \\sim 1$.  This suggests a new 2-component model for galaxy formation that predates the appearance of the red sequence.  \\item We propose two interpretations for this bimodal distribution, focusing on the ``dip'' in the blue mass function at $\\sim$10$^{10} \\Msun$.  If the gas fraction increases at lower stellar masses, galaxies with $M_{\\rm baryon} \\sim 10^{10} \\Msun$ would shift to lower stellar masses, creating the observed dip.  This would indicate a change in SFR efficiency, perhaps arising from the influence of supernovae feedback which likely sets in below scales of $\\sim$10$^{10} \\Msun$. In this picture, the baryonic mass function should not show a dip. Using published (cold) gas fractions as a function of stellar mass, we show that cold gas alone is not sufficient to eliminate the dip in the baryonic mass function, but that the addition of the hard to detect warm gas could potentially flatten out the baryonic mass function considerably.  \\item Alternatively, the dip could be created by an enhancement of the galaxy assembly rate at $\\sim$10$^{11} \\Msun$, a phenomenon that naturally arises if the baryon fraction peaks at $M_{\\rm halo} \\sim 10^{12} \\Msun$. In this scenario we would identify the galaxies occupying the bump around $\\Mstar$ with central galaxies and the increasing fraction of satellite galaxies at lower mass with the second, fainter, component of the mass function and in particular the steep faint-end slope.  \\item The low-mass end of the blue and total mass functions exhibit a steeper slope than has been detected in previous work.  This can be interpreted as a {\\em steepening} slope, one that may increasingly approach the halo mass function value of -2.  \\item While the dip feature is apparent in the total mass function at all redshifts, it appears to shift from the blue to red population, likely as a result of transforming high-mass blue galaxies into red ones.  \\item At the same time, we detect a drastic upturn in the number of low mass red galaxies.  Their increase with time seems to reflect a decrease in the number blue systems and so we tentatively associate them with satellite (dwarf spheroidal) galaxies that have undergone quenching due to environmental processes.  \\end{itemize}  While the broad dynamic range of the COSMOS data set allows us to begin to characterize some of the more subtle features of the stellar mass function, the single COSMOS field is still limited by cosmic variance. Future work over several fields will verify the trends here and shed light on possible interpretations by constraining how these features evolve with redshift."
    },
    "0910/0910.5666_arXiv.txt": {
        "abstract": "We present the first X-ray analysis of the diffuse hot ionized gas and the point sources in IC131, after NGC604 the second most X-ray luminous giant \\hii\\ region in \\m. The X-ray emission is detected only in the south eastern part of IC131 (named IC131-se) and is limited to an elliptical region of $\\sim$200 pc in extent. This region appears to be confined towards the west by a hemispherical shell of warm ionized gas and only fills about half that volume. Although the corresponding X-ray spectrum has 1215 counts, it cannot conclusively be told whether the extended X-ray emission is thermal, non-thermal, or a combination of both. A thermal plasma model of $kT_e$\\,=\\,4.3\\,keV or a single power law of $\\Gamma$\\,$\\simeq$\\,2.1 fit the spectrum equally well. If the spectrum is purely thermal (non-thermal), the total unabsorbed X-ray luminosity in the 0.35\\,--\\,8\\,keV energy band amounts to $L_X=6.8\\ (8.7)\\times 10^{35}$\\,\\ergs. Among other known \\hii\\ regions IC131-se seems to be extreme regarding the combination of its large extent of the X-ray plasma, the lack of massive O stars, its unusually high electron temperature (if thermal), and the large fraction of $L_X$ emitted above 2\\,keV ($\\sim$40\\,--53\\%). A thermal plasma of $\\sim$4\\,keV poses serious challenges to theoretical models, as it is not clear how high electron temperatures can be produced in \\hii\\ regions in view of mass-proportional and collisionless heating. If the gas is non-thermal or has non-thermal contributions, synchrotron emission would clearly dominate over inverse Compton emission. It is not clear if the same mechanisms which create non-thermal X-rays or accelerate CRs in SNRs can be applied to much larger scales of 200\\,pc. In both cases the existing theoretical models for giant \\hii\\ regions and superbubbles do not explain the hardness and extent of the X-ray emission in IC131-se. We also detect a variable source candidate in IC131. It seems that this object is a high mass X-ray binary whose optical counterpart is a B2-type star with a mass of $\\sim$9$M_{\\odot}$. ",
        "introduction": "This study is part of the \\cxo\\ ACIS Survey of M33 \\citep[\\chase,][]{plu08} and presents the first deep, high resolution X-ray images of \\object{IC131}. Within \\chase's sensitivity limit of $\\sim$$2\\times 10^{34}$\\, erg\\,s$^{-1}$ (0.35\\,--\\,8\\,keV) IC131 is, along with \\object{NGC604} \\citep{tull08}, the only giant \\hii\\ region (GHR) in \\object{\\m}\\ which shows significant diffuse extended X-ray emission.  Because of their large sizes, masses, stellar contents, and luminosities, GHRs \\citep[cf.][]{ken84} provide excellent laboratories for studying the relation between stellar feedback and the evolution of the individual components of the ISM inside them. GHRs like 30\\,Dor \\citep{wang99,town06}, N51D \\citep{coop04}, or NGC604 \\citep{tull08} are strong X-ray emitters thanks to their massive O star population. The heavy mass loss of these stars can produce strong colliding winds which, together with contributions from SNe and SNRs, shock-heat the gas and forces it to emit thermal X-rays at temperatures around $kT_e=0.6\\pm0.2$\\,keV. In some cases the X-ray spectra of \\hii\\ regions in the Milky Way and the LMC require, in addition to a thermal component, also a non-thermal component \\citep[e.g.,][]{wolk02,bam04,coop04,smith04,muno06,mad09}. The non-thermal radiation could be explained by synchrotron emission, inverse Compton scattering, or by particle acceleration in shock regions \\citep{par04,arsch08}. In all cases, the X-ray emission could be linked to massive stars or their successors. The issue one faces with GHRs is to find a mechanism which can ionize the gas on spatial scales which can easily exceed that of normal \\hii\\ regions by one to two orders of magnitude and is consistent with the stellar energy feedback, X-ray luminosity, gas mass, and other predictable quantities.  Although IC131 has been observed in numerous M33 surveys across almost all wavelengths \\citep[e.g.,][]{rosa84,land92,long96,hippe03,pietsch04,mass06,tab07,plu08}, little is known about the stellar population, its age, and the different components of the ISM. The \\chase\\ data set, supplemented by other observations at optical and infrared wavelengths, is the first data set that has the combination of S/N and spatial resolution to permit a detailed study of the X-ray emitting gas in IC131. Therefore, the primary focus of this work is to study the hot ionized medium (HIM) by constraining basic plasma parameters, to determine the likely ionization mechanism of the hot gas, and to compare properties of the gas to other star forming regions to learn more about the evolution of the gas in GHRs. The second purpose of this study focuses on the analysis of the point source population in IC131.  In the following, we give an overview of the observations used in this study and the data reduction steps that have been applied (Section \\ref{sec2}). We then present high resolution (2$\\arcsec$) images of the diffuse X-ray emission in different energy bands and compare these to multi-wavelength data to constrain the morphology of IC131 (Section \\ref{spec_img}). A spectral analysis is carried out in Section \\ref{spec_res} while Section \\ref{spec_dis} discusses the possible origin and ionization mechanism of the HIM. Section \\ref{point_dis} deals with the analysis of the X-ray emitting point sources detected in IC131 and Section \\ref{con} gives a concluding summary of the most important results. ",
        "conclusions": "\\label{con} Compared to known superbubbles and \\hii\\ regions IC131-se appears to be peculiar with respect to its lopsided X-ray morphology, the large linear extent of the X-ray emission, the lack of massive O stars, its high electron temperature (in case the X-ray gas is thermal), and large fraction of hard X-rays. The X-ray spectrum of the extended emission in IC131-se can be equally well fitted by an absorbed power law ($\\Gamma \\simeq 2.1$) or an absorbed thermal plasma model ($kT_e\\simeq 4.3$\\,keV). These models predict a total (0.35\\,--\\,8.0\\,keV) unabsorbed X-ray luminosity of about $9\\times10^{35}$\\ergs\\ and $7\\times10^{35}$\\ergs, respectively, with 39\\% and 53\\% of the luminosity being emitted above 2\\,keV. Apparently IC131 possesses not only the hardest X-ray spectrum among the known GHRs, but, in case the gas is thermal, also the hottest X-ray plasma.  Besides the facts that thermal X-ray emission is expected from GHRs and that there might be weak emission lines below 1\\,keV, the disagreement between the number of SNe progenitors obtained from the IMF (9\\,--\\,21) and that derived from the internal thermal energy (118$\\pm$59) argues against a pure thermal origin of the gas. We estimate a low electron density of $n_e\\simeq 0.05$\\,cm$^{-3}$ (assuming $f_X =0.9$) and an X-ray gas mass of about 3300\\,$M_{\\odot}$. The age of the stellar population inside the bubble seems to range from 8 to 77\\,Myr and given a cooling timescale of about 1\\,Gyr, the gas did not have time to cool significantly. If the gas is thermal, the standard bubble models require shocks from SNRs and SNe to heat the gas. The fundamental problem which argues against a thermal plasma is that we are not aware of a mechanism for GHRs which can produce electron temperatures anywhere close to $kT_e=4$\\,keV.  Non-thermal X-ray emission is also known to originate from SNRs, GHRs or superbubbles and the spectral power law fit is statistically as good as the one for a thermal plasma. Synchrotron emission as opposed to inverse Compton scattering is a plausible mechanism for the assumed magnetic field strength of $B=20\\mu G$ and the total energy of the relativistic synchrotron-emitting electron population is about 14\\% to 30\\% of the total energy provided by SNe. If a non-thermal component exists, synchrotron losses clearly dominate over losses from inverse Compton scattering. For a purely non-thermal origin, however, it remains to be investigated what kinds of mechanisms are able to produce non-thermal X-rays or accelerate particles to relativistic energies on scales of 200\\,pc. To allow for a more quantitative analysis of the non-thermal X-ray emission, models are required which predict the fraction of non-thermal radiation from GHRs. A combined thermal and non-thermal model seems to be generally appropriate, too, irrespective of the unreasonably high column density and X-ray luminosity derived from those fits.  Unfortunately, with the present data the nature of the extended X-ray emission cannot be conclusively determined. Clearly, IC131-se is an important object and challenges the standard bubble model as well as our understanding of CR acceleration in superbubbles. It is remarkable and remains to be understood why objects similar to IC131-se have not been detected before. Future investigations would greatly profit from more sensitive and deeper X-ray observations (as for example provided by IXO), high-resolution optical spectrophotometric data, and models which can explain electron temperatures $>$\\,1\\,keV and make predictions on the non-thermal X-ray emission in GHRs. The deeper X-ray observations could establish the presence of X-ray emission lines and settle the question on the nature of the emission, while a detailed stellar population analysis could provide a more accurate IMF. Clearly, NGC604 and IC131-se seem to be in completely different evolutionary stages. All O-type stars in IC131-se seem to have exploded as SNe, whereas the western part of NGC604 awaits the first SNe to occur.  We detect only one X-ray point source in IC131. This source, CXO\\,J013315.10+304453.0 (FL073), appears to be time variable and is possibly a HMXRB with an optical counterpart which could be a B2V star with a mass of about 9M$_{\\odot}$."
    },
    "0910/0910.2270_arXiv.txt": {
        "abstract": "Spectral line and continuum observations of the ionized and molecular gas in G20.08-0.14 N explore the dynamics of accretion over a range of spatial scales in this massive star-forming region. Very Large Array observations of NH$_3$ at $4\"$ angular resolution show a large-scale (0.5 pc) molecular accretion flow around and into a star cluster with three small, bright \\HII regions. Higher resolution ($0.4''$) observations with the Submillimeter Array in hot core molecules (CH$_3$CN, OCS, and SO$_2$) and the VLA in NH$_3$, show that the two brightest and smallest \\HII regions are themselves surrounded by smaller scale (0.05 pc) accretion flows. The axes of rotation of the large and small scale flows are aligned, and the timescale for the contraction of the cloud is short enough, 0.1 Myr, for the large-scale accretion flow to deliver significant mass to the smaller scales within the star formation timescale. The flow structure appears to be continuous and hierarchical from larger to smaller scales.   Millimeter radio recombination line (RRL) observations at $0.4\\arcsec$ angular resolution indicate rotation and outflow of the ionized gas within the brightest \\HII region (A). The broad recombination lines and a continuum spectral energy distribution (SED) that rises continuously from cm to mm wavelengths, are both characteristic of the class of \\HII regions known as ``broad recombination line objects''. The SED indicates a density gradient inside this \\HII region, and the RRLs suggest supersonic flows. These observations are consistent with photoevaporation of the inner part of the rotationally flattened molecular accretion flow.  We also report the serendipitous detection of a new NH$_3$ (3,3) maser. ",
        "introduction": "\\label{intro}  Massive star-forming regions (MSFRs) with O stars are usually identified by a group of hypercompact (HC) \\HII or ultracompact (UC) \\HII regions found together, deeply embedded in a dense molecular cloud \\citep[][]{GL09,Church02,Hoar07}. That several \\HII regions are typically found within each star-forming region indicates that massive stars form together in small clusters. Furthermore, the infrared luminosity and radio continuum brightness of the individual \\HII regions suggest that some of them may themselves contain more than one massive star. Thus, the spatial structure of massive star-forming regions is clustered and hierarchical: the star-forming regions contain a number of separate HC and UC \\HII regions, each of which may in turn contain a few massive stars.  Low angular resolution, single-dish, molecular line surveys of MSFRs show evidence for large scale contraction of the embedding molecular clouds \\citep{WE03,KW08}. Higher angular resolution observations of some of these regions identify velocity gradients consistent with rotation and inflow. In addition to the accretion flows seen on the large-scale ($\\sim 0.3-1$ pc) of the embedding molecular cloud (G10.6-0.4: Ho \\& Haschick 1986, Keto et al. 1987a, Keto 1990; G29.96-0.02: Olmi et al. 2003), accretion flows are also seen on smaller ($\\leq 0.1$ pc) scales around individual HC and UC \\HII regions (G10.6-0.4: Keto et al. 1988, Sollins et al. 2005a; W3(OH): Keto et al. 1987b, Keto et al. 1995; W51e2: Zhang \\& Ho 1997, Young et al. 1998; G28.20-0.05: Sollins et al. 2005b; G24.78+0.08: Beltr\\'an et al. 2004, 2005, 2006, Galv\\'an-Madrid et al. 2008; G29.96-0.02: Beuther et al. 2007).  It is unclear how the flows on different length scales are related. In the case of G10.6-0.4, the cluster-scale accretion flow can be traced down from the largest cloud scale to the small scale of the brightest \\HII region, but it is not known whether this holds for other objects. For example, in a survey of MSFRs, selected on the basis of IRAS colors and specifically excluding those with \\HII regions, multiple bipolar molecular outflows (implying the presence of accretion flows) are seen in random orientations \\citep{Beu02a,Beu02,Beu03}. The different orientations of these smaller-scale flows suggest separate, individual centers of collapse. This comparison raises the question whether a large-scale coherent flow is required for the formation of the most massive stars, O stars ($M_\\star > 20$ $\\Msun$) capable of producing bright \\HII regions, whereas B stars require only smaller scale flows.  It is also unclear what happens in an accretion flow when the inflowing molecular gas reaches the boundary of an embedded \\HII region. Previous observations suggest that the \\HII regions in an MSFR that are surrounded by accretion flows, may be best understood as deriving from the continuous ionization of the accretion flow \\citep{Keto02, Keto03, Keto07}, rather than as a dynamically separate expanding bubble of ionized gas within the flow. Part of the ionized gas may continue to the central star or stars and part escapes off the rotationally flattened accretion flow as a photoevaporative outflow perpendicular to the plane of rotation \\citep{Hol94, YW96, Liz96, John98, Lugo04}. The outflow is accelerated to supersonic speeds by the density gradient maintained by the stellar gravity \\citep{Keto07}. Because the extent of an ionized outflow is generally larger than the region of ionized inflow, in most cases the outflow should be detected more easily than the inflow. \\HII regions classified as ``broad recombination line objects\" (BRLO) \\citep{JaffMP99,Sewi04} show steep density gradients and supersonic flows \\citep{KZK08}, consistent with photoevaporation and acceleration. It is not known whether all BRLO are associated with accretion. If the accretion surrounding an O star cluster is continuous from the largest to the smallest scales, this must be the case.   There are only a handful of radio recombination line (RRL) observations that spatially resolve the ionized flow within an HC \\HII region. Velocity gradients consistent with outflow and rotation in the ionized gas have been previously reported for W3(OH) \\citep{Keto95}, W51e2 \\citep{KK08}, and G28.20-0.05 \\citep{Sewi08}. Observations of the very massive and spatially large G10.6-0.4 \\HII region made at the VLA in the highest possible angular resolution are able to map the inflowing ionized gas \\citep{KW06}.   In order to study the accretion dynamics over a range of scales in a MSFR, from the cluster scale down to the scale of individual HC \\HII regions and within the ionized gas, we set up a program of radio frequency molecular line, recombination line, and continuum observations at two telescopes and with several different angular resolutions. For this study we chose the massive star formation region G20.08-0.14 North (hereafter G20.08N), identified by three UC and HC \\HII regions detected in the cm continuum by \\cite{WC89}. The total luminosity of the region is $L \\sim 6.6 \\times 10^5$ $L_\\odot$ for a distance of 12.3 kpc.\\footnote{ Both near and far kinematic distances have been reported for G20.08N. The near value given by \\cite{Dow80} ($d \\approx 4.1$ kpc) is the most commonly quoted in the previous literature. In contrast, \\cite{Fish03} and \\cite{AnBan09} report that this region is at the far kinematic distance ($d \\approx 12.3$ kpc). We will assume the far distance throughout the rest of the paper. For reference, a scale of $0.5\\arcsec$ corresponds to $\\approx 6000$ AU (0.03 pc). The total luminosity of the region was estimated to be $L \\sim 7.3 \\times 10^4$ $L_\\odot$ assuming the near kinematic distance \\citep{WC89}. Correcting for the location at the far distance, the luminosity is $L \\sim 6.6 \\times 10^5$ $L_\\odot$.}   Previous observations suggest accretion in the G20.08N cluster. Molecular-line observations show dense gas embedding the \\HII regions \\citep{Tur79,Plume92}. Molecular masers, generally associated with ongoing massive-star formation, have been detected in a number of studies (OH: Ho et al. 1983; H$_2$O: Hofner \\& Churchwell 1996; and CH$_3$OH: Walsh et al. 1998). \\citet{KW07,KW08} observed large-scale inward motions consistent with an overall contraction of the embedding molecular cloud. Those authors also observed SiO line profiles suggestive of massive molecular outflows, further evidence for accretion and star formation. The recombination line spectra show broad lines \\citep{Gar85, Sewi04}, presumably due to large, organized motions in the ionized gas. However, the previous observations do not have the angular resolution and the range of spatial scales needed to confirm the presence of accretion flows and study them in detail.   In this paper we report on several observations of G20.08N and discuss our findings. We confirm active accretion within the cluster. Furthermore, we find that the parsec-scale accretion flow fragments into smaller flows around the individual HC \\HII regions, and  that the gas probably flows from the largest scale down to the smallest scale. This continuous and hierarchical accretion may be necessary to supply enough mass to the small-scale flows to form O-type stars, in contrast to low- and intermediate-mass star-forming regions with stars no more massive than $\\sim 20$ $\\Msun$, where isolated accretion flows around individual protostars may be sufficient. ",
        "conclusions": "\\label{conclu}  We report radio and mm observations of the molecular and ionized gas toward the O-star cluster G20.08N, made with an angular resolution from $\\sim 0.1$ pc to $\\sim 0.01$ pc.  Our main findings can be summarized as follows:  \\begin{enumerate}  \\item We find a large-scale ($\\sim 0.5$ pc) accretion flow around and into a star cluster with several O-type stars, identified by one UC and two HC \\HII regions. This flow is rotating and infalling towards its center. The two HC \\HII regions are surrounded by smaller accretion flows ($\\sim 0.05$ pc), each of them with the signature of infall too. The brightest (toward \\HII region A) is detected in mm emission lines, and rotates in concordance with the large-scale flow.  \\item The similar orientations of the flows at small and large scales, as well as their dynamical timescales ($\\sim 10^4$ yrs and $\\sim 10^5$ yrs respectively), and masses ($\\sim 10~\\Msun$ and $\\sim 10^3~\\Msun$ respectively), suggest that, if O stars are forming in G20.08N (as it is observed), then the smaller scales ought to be resupplied from the larger scales. The same result has been found in recent numerical simulations of massive star formation in clusters.  \\item The brightest HC \\HII region (A) has a rising SED from cm to mm wavelengths and broad hydrogen recombination lines. Both characteristics suggest density gradients and supersonic flows inside the \\HII region. A tentative velocity gradient is detected in the recombination line emission of this source, suggesting rotation and outflow in the ionized gas at the innermost scales. \\HII region A can be interpreted as the inner part of the surrounding molecular accretion flow, with the observed ionization being produced by photoevaporation.   \\end{enumerate}"
    },
    "0910/0910.0275_arXiv.txt": {
        "abstract": "We present blue optical spectra of 92 members of $h$ and $\\chi$ Per obtained with the WIYN telescope at Kitt Peak National Observatory. From these spectra, several stellar parameters were measured for the B type stars, including $V$ sin $i$, $T_{\\rm eff}$, log $g_{\\rm polar}$, $M_{\\star}$, and $R_{\\star}$. Str\\\"{o}mgren photometry was used to measure $T_{\\rm eff}$  and log $g_{\\rm polar}$ for the Be stars. We also analyze photometric data of cluster members and discuss the near-to-mid IR excesses of Be stars. ",
        "introduction": "NGC 869  and NGC 884 ($h$ and $\\chi$ Persei, respectively) are a well known double cluster rich in massive B-type stars, and have been the focus of many studies over the years. Recent studies show that NGC 869 and NGC 884 have nearly identical ages of $\\sim$ 13--14 Myr, common distance moduli of dM $\\sim$ 11.85, and common reddenings of E(B-V) $\\sim$ 0.5--0.55 (\\cite[Currie et al.\\ 2009]{Currie_etal09}, \\cite[Slesnick et al.\\ 2002]{Slesnick_etal02}, \\cite[Bragg \\& Kenyon 2005]{Bragg_etal05}).  \\cite[Currie et al.\\  (2008; hereafter C08)]{Currie_etal08} identified two populations of NGC 869 and NGC 884 stars with detected Spitzer MIPS-24 $\\mu$m excess emission: 20 A and F-type stars with luminous debris disk emission and 57 brighter, earlier stars with weaker excess emission. They identify most of the latter group as candidate Be stars. However, only 21 were previously listed as Be stars (eg. \\cite[Bragg \\& Kenyon 2002]{Bragg_etal02}, \\cite[Slesnick et al.\\ 2002]{Slesnick_etal02}).  In this study, we analyze blue optical spectra of 92 early-type cluster members, including 16 candidate Be stars from C08, and investigate their near-to-mid infrared (IR) excesses. With continued monitoring of these stars in the both the optical and IR regimes, we hope to explore these excesses as a reasonable means for identifying potential Be stars within clusters, as well as to investigate the transient natures of the disks surrounding the known Be stars in NGC 869 and NGC 884.  \\vspace{-0.1in} ",
        "conclusions": "We have measured the physical parameters of 77 B-type stars and 15 Be stars in NGC 896 and NGC 884. Sixteen Be candidates from C08 are present in our sample or that of HG06. Of these 16 Be candidates, 3 stars show no evidence of emission in our optical data and are likely transient Be stars. Ten of these Be candidates do show emission in our spectra. Those Be candidates without emission in our spectra should be monitored in the future to further investigate their transient nature.  In the future, IRAC 3.6-5.8 $\\mu$m data will be combined with the optical and IR fluxes used here to investigate the observed SEDs. We will fit the new SEDs using modern flux models rather than blackbody curves. Modifications accounting for variable reddening throughout the clusters will also be made. These new SED fits can then be used to model the Be disk sizes and temperatures."
    },
    "0910/0910.2877_arXiv.txt": {
        "abstract": "Current X-ray observatories make it possible to follow the evolution of transient and variable X-ray binaries across a broad range in luminosity and source behavior.  In such studies, it can be unclear whether evolution in the low energy portion of the spectrum should be attributed to evolution in the source, or instead to evolution in neutral photoelectric absorption.  Dispersive spectrometers make it possible to address this problem.  We have analyzed a small but diverse set of X-ray binaries observed with the {\\it Chandra} High Energy Transmission Grating Spectrometer across a range in luminosity and different spectral states.  The column density in individual photoelectric absorption edges remains constant with luminosity, both within and across source spectral states.  This finding suggests that absorption in the interstellar medium strongly dominates the neutral column density observed in spectra of X-ray binaries.  Consequently, evolution in the low energy spectrum of X-ray binaries should properly be attributed to evolution in the source spectrum.  We discuss our results in the context of X-ray binary spectroscopy with current and future X-ray missions. ",
        "introduction": "X-ray binaries show a remarkable variety of extreme phenomena, including rapid variability in many wavelength bands, the production of relativistic jets and hot winds, and outbursts that span factors of $10^{6}$ in luminosity (for a review of black hole transients, see Remillard \\& McClintock 2006).  Monitoring observations with the {\\it Rossi X-ray Timing Explorer} and {\\it Swift} enable the evolution of X-ray spectral and timing properties of X-ray binaries to be traced in order to understand the accretion process (see, e.g., Rykoff et al.\\ 2007).  Model degeneracies that may occur in fitting a single spectrum can often be resolved by examining how multiple spectra evolve with time.  Major changes within the outbursts of transient and persistent X-ray binaries are typically linked with changes in the nature of the accretion disk (see, e.g., Esin, McClintock, \\& Narayan 1997; Homan et al.\\ 2001; also see Miller et al.\\ 2006a).  Tracking accretion disk evolution is therefore a primary goal of monitoring observations.  The disk is directly visible in soft X-rays in black hole X-ray binaries, and likely also in neutron star binaries (e.g. Cackett et al.\\ 2008a).  However, the K shell photoelectric absorption edges from most abundant elements fall below 2~keV, as do Fe L shell edges.  This absorption can complicate efforts to study disk properties, especially when using a spectrometer with limited energy range and resolution.  Prominent theoretical treatments of accretion flows predict modest winds (e.g. Begelman, McKee, \\& Shields 1983) and/or outflows with a high ionization parameter (e.g. Narayan \\& Raymond 1999). Observationally, even the most extreme winds detected in X-ray binaries do not appear to generate enough column density to add appreciably to the neutral absorption column (Miller et al.\\ 2006). Nevertheless, it is not uncommon to allow the equivalent neutral hydrogen column density to vary in modeling multiple spectra of an X-ray binary (see, e.g. Brocksopp et al.\\ 2006; Caballero-Garcia et al.\\ 2009; Cabanac et al.\\ 2009).  Even in cases where the neutral column density is held constant or jointly determined, this method is adopted somewhat arbitrarily, without specific reference to an observational or theoretical basis (see, e.g., Sobczak et al.\\ 2000; Miller et al.\\ 2001; Nowak, Wilms, \\& Dove 2002; Park et al.\\ 2004; Gierlinski \\& Done 2004; Kalemci et al.\\ 2005; Rykoff et al.\\ 2007).  Dispersive X-ray spectroscopy can separate variations in absorption from variations in the source spectrum.  Even the resolution afforded by CCD spectrometers blends many weak edges into the continuum, and the optical depth and wavelength of stronger edges can be difficult to measure precisely.  For instance, in fits to CCD spectra, the multiple Fe L edges can be modeled satisfactorily in terms of a single step function.  In constrast, {\\it Chandra} High Energy Transmission Grating Spectrometer observations of X-ray binaries have resolved the individual Fe L edges and their detailed properties (see, e.g., Schulz et al. 2002, Juett et al.\\ 2006).  Beyond simple edges, dispersive X-ray spectroscopy has clearly revealed the neutral O I 1s-2p resonance absorption line that was first glimpsed with the crystal spectrometer aboard {\\it Einstein} (Schattenberg \\& Canizares 1986, Paerels et al.\\ 2001).  In this Letter, we extend recent work using high resolution X-ray spectroscopy by examining the evolution of neutral absorption in a sample of X-ray binaries.  We find no evidence for variable absorption, consistent with an interstellar origin for neutral absorption. ",
        "conclusions": "To better understand the evolution of the low energy spectrum of X-ray binaries, we made fits to individual photoelectric absorption edges in high resolution X-ray spectra of selected sources.  The column density measured in individual edges is not observed to vary across different spectral states, nor over a broad range in luminosity (see Table 1 and Figure 2).  This suggests that gas from X-ray binaries is not typically an important source of the neutral absorption observed in the spectra of these systems.  Rather, neutral absorption must be dominated by the ISM.  A similar conclusion was reached by Juett, Schulz, \\& Chakrabarty (2004) based on upper limits on the velocity dispersion of the ISM as measured through absorption lines.  Evolution in the low energy spectrum of typical X-ray binaries, then, is best attributed to evolution in the source continuum.  Neutral absorption in X-ray spectra is often fit by a single model that parameterizes the accumulated absorption from individual edges as an equivalent neutral hydrogen column density.  Values obtained from high resolution spectra are likely to give the best measure of a true equivalent total column density.  In practice, instrumental problems such as internal scattering, carbon build-up from optical blocking filters, and gain drift could prevent the adoption of a gratings-derived value for the neutral column.  In such cases, our results suggest that a single value of the equivalent neutral hydrogen column density should be used to fit multiple spectra from monitoring observations of a given source with a given detector.  The outflows that are observed in X-ray binaries are highly ionized -- dominated by He-like and H-like charge states (see, e.g., Lee et al.\\ 2002, Miller et al.\\ 2004, Miller et al.\\ 2006b, Schulz et al.\\ 2008).  An especially dense wind was observed in GRO J1655$-$40, and even in that case the ionized columns observed are insufficient to create strong absorption edges that could be mistaken for additional neutral absorption (Miller et al.\\ 2006c, 2008). Moreover, in sources such as H 1743$-$322, GRO J1655$-$40, and GRS 1915$+$105, a paradigm is emerging wherein ionized winds are active in soft, disk-dominated states, but absent in spectrally hard states like those that typically hold when sources accrete at a low fraction of their Eddington limit (Miller et al.\\ 2006b, 2006c, Miller et al.\\ 2008, Neilsen \\& Lee 2009).  In Figure 2, it is clear that any variation in ionized winds across states does not affect the column density measured in neutral edges.  Our results are based on spectra which only reach down to approximately $0.01~{\\rm L}_{\\rm Edd.}$.  Sensitive high-resolution X-ray spectra that would permit strong constraints on absorption variability have not yet been obtained from sources at lower luminosity.  However, theoretical arguments again point to ionized outflows that would contribute little to a neutral column.  Winds from advection-dominated accretion flows are expected to be extremely hot (since advective flows are very hot), and line spectra should be dominated by He-like and H-like charge states (e.g. Narayan \\& Raymond 1999).  Recent observations of the stellar-mass black hole V404 Cyg, which accretes at about $10^{-5}~ {\\rm L}_{\\rm Edd.}$, are able to rule-out the winds predicted by some advective models (Bradley et al.\\ 2007).  This further suggests that ionized outflows are not likely to contribute significantly to line-of-sight absorption, even as sources approach quiescence.  There are at least two classes of sources where our results may not hold in all circumstances.  Neutron star X-ray binaries known as ``dippers'' are viewed at high inclinations, and material in the outer disk can block emission from the central engine (see, e.g., Diaz Trigo et al.\\ 2006).  Within such flux dips, the observed neutral absorption column may vary due to the changing geometry within the binary.  Similarly, winds from massive stars are known to be clumpy and to sometimes cause flux dips; some of these dips may also cause variations in the observed equivalent neutral hydrogen column density (e.g. Balucinska-Church et al.\\ 2000).  \\vspace{0.1in}  \\noindent We thank Joern Wilms for creating the ``tbnew'' model used in this work, and for guidance on its implementation.  We acknowledge helpful comments from the anonymous referee that improved this paper. J.M.M. gratefully acknowledges funding from the {\\it Chandra} Guest Observer program.  EMC gratefully acknowledges support provided by NASA through the Chandra Fellowship Program, grant number PF8-90052. RCR acknowledges STFC for financial support."
    },
    "0910/0910.0043_arXiv.txt": {
        "abstract": "{} {We are studying the dust properties of four low metallicity galaxies by modelling their spectral energy distributions. This modelling enables us to constrain the dust properties such as the mass, the temperature or the composition to characterise the global ISM properties in dwarf galaxies. } {We present 870 \\mic\\ images of four low metallicity galaxies (NGC~1705, Haro~11, Mrk~1089 and UM~311) observed with the Large \\APEX\\ BOlometer CAmera (LABOCA) on the Atacama Pathfinder EXperiment (\\APEX) telescope. We model their spectral energy distributions combining the submm observations of \\lab, \\twomass, \\iras, \\spitz\\ photometric data and the IRS data for Haro~11.} {We find that the PAH mass abundance is very low in these galaxies, 5 to 50 times lower than the PAH mass fraction of our Galaxy. We also find that a significant mass of dust is revealed when using submm constraints compared to that measured with only mid-IR to far-IR observations extending only to 160 \\mic. For NGC~1705 and Haro~11, an excess in submillimeter wavelengths is detected when we use our standard dust SED model. We rerun our SED procedure adding a cold dust component (10 K) to better describe the high 870 \\mic\\ flux derived from LABOCA observations, which significantly improves the fit. We find that at least 70 $\\%$ of the dust mass of these two galaxies can reside in a cold dust component. We also show that the subsequent dust-to-gas mass ratios, considering HI and CO observations, can be strikingly high for Haro~11 in comparison with what is usually expected for these low-metallicity environments. Furthermore,  we derive the star formation rate of our galaxies and compare them to the Schmidt law. Haro~11 falls anomalously far from the Schmidt relation. These results may suggest that a reservoir of hidden gas could be present in molecular form not traced by the current CO observations. While there can be a significant cold dust mass found in Haro~11, the SED peaks at exceptionally short wavelengths (36 \\mic), also highlighting the importance of the much warmer dust component heated by the massive star clusters in Haro~11. We also derive the total IR luminosities derived from our models and compare them with relations that derive this luminosity from \\spitz\\ bands. We find that the Draine $\\&$ Li (2007) formula compares well to our direct IR determinations. } {} ",
        "introduction": "The understanding of the evolution of a galaxy requires knowledge of the roles of the different actors controlling the evolution of the Interstellar Medium (ISM) and the subsequent feedback on star formation activity. Despite its low fraction of the total mass of a galaxy (less than 1$\\%$), dust plays a prominent role in the heating and cooling of the ISM and thus tightly influences the overall physics of a galaxy. Since dust absorbs the stellar radiation and reemits it in a wide range of wavelengths, the star formation rate (SFR) as well as other fundamental parameters of a galaxy, such as its age, can be indirectly studied through the dust emission itself. {\\revisedbis The Spectral Energy Distribution (SED) of a galaxy is its spectral footprint from which we can study the physical processes taking place in the galaxy since it synthesises the contribution of all its components to the emission of the galaxy. Using this tool, we can peer into the window of the integrated history of the galaxy and disentangle the various physical actors (stars, HII regions, molecular clouds) and processes (stellar radiation, dust emission) involved~\\citep[][see also $\\S$ 5 of this paper]{Draine2007,Galliano_Dwek_Chanial_2008}.}  While dust hinders the interpretation of ultraviolet (UV) and optical wavelengths, in the Mid Infrared (MIR), Far Infrared (FIR) and submillimetre (submm) wavelengths, dust emission and absorption properties expose different physical environments, from the most vigorous star formation and AGN activity \\citep[e.g.][]{Gordon1995, Wu2007} to the more quiescent diffuse media \\citep{Bernard1996,Arendt1998}. Many processes linked to star formation such as stellar winds \\citep{Hoefner2009}, supernovae shocks, photodestruction by high-mass stars etc. can also affect the spatial distribution and the local properties and abundance of the different dust components of a galaxy such as Polycyclic Aromatic Hydrocarbons \\citep[PAHs,][]{OHalloran2006}, amorphous carbon grains, silicates or composite grains, manifesting themselves in the MIR to submm wavelengths.  Studying the interplay between galaxy properties and metal enrichment is crucial to understand galaxy evolution. The metallicity of a galaxy is deeply linked with the dust properties of the ISM and its substructures such as HII regions and molecular clouds, but just how it affects the ISM is currently poorly known. Dwarf galaxies in the Local Universe, are metal-poor galaxies, and are thus convenient laboratories to study the effects of metallicity (Z) on the gas and dust. They exhibit a wide variety of physical conditions, and their star formation properties and ISM represent the closest analogs to proto-galaxies of the early universe. Indeed, dwarf galaxies are small and may compared to high redshift galaxies which also present lower metallicities \\citep{Lara_Lopez_2009}. They are also considered to be the building blocks of much larger and more metal-rich galaxies \\citep[Review by][]{Tosi2003}. They also show analogies with Gamma Ray Bursts (GRB) hosts whose ISM usually exhibit moderate chemical enrichment with a median metallicity of 1/10 \\zsun \\citep{Chen2009}.  They finally show evidence for older stellar populations than their metallicity suggests \\citep[e.g.][]{Aloisi1998}, posing enigmatic issues for galaxy evolution models. Many studies have been carried out to grasp this apparent paradox. \\citet{Lisenfeld_Ferrara_1998}  confirmed that the dependence of the dust-to-gas mass ratio (D/G) in low metallicity galaxies was a function of metallicity using \\iras\\ observations.  \\citet{James2002}, \\citet{Walter2007} and \\citet{Hirashita2008} concluded likewise using \\jcmt/SCUBA submm, \\spitz\\ MIR/FIR and \\akari\\ (FIR)  observations. Finally, \\citet{Galliano_Dwek_Chanial_2008} observed some systematic deviations between dust abundances of very low metallicity systems and what is expected for supernova-condensed dust. At MIR wavelengths, low metallicity systems also show prominent differences in the dust properties compared to the more metal-rich systems. For example, PAH features are strikingly diminished as metallicities drop \\citep[e.g.][]{Madden2005, Engelbracht2005, Wu2007, Engelbracht2008} compared to metal-rich galaxies, in spite of the role the smallest grains play in the energy balance of galaxies \\citep{Rubin2009}.  Some studies suggest that PAH emission depends on the hardness or strengh of the illuminating radiation field  \\citep{Madden2006, Engelbracht2008, Gordon2008, Bendo2008}. The consequence of lowering the metallicity of a galaxy is the decrease in dust opacity resulting in harder and stronger radiation fields. The dearth of PAHs in low metallicity galaxies has also been explained by the destructive effects of supernovae \\citep{OHalloran2006,OHalloran2008} or by the delayed injection of PAHs by AGB stars \\citep{Galliano_Dwek_Chanial_2008}.  Broad wavelength coverage of the MIR to submm regime is imperative to constrain the modelling of the observed SEDs, leading to a better comprehension of the dust properties of galaxies. Since \\spitz\\  only observes dust emission at wavelengths shorter than 160 \\mic, submm data are necessary not only to enlarge the wavelength coverage at longer wavelengths to verify the dust models but also because the potential reservoir of cold grains ($\\le$ 15K), which contribute to this submm flux, may account for a significant amount of mass. Only a handful of galaxies of the Local Universe have been studied using submm ground-based instruments (e.g. \\jcmt/SCUBA). When submm observations of dwarf or late-type galaxies are studied, an excess in the dust SEDs is often found in the mm/submm domain \\citep{ Lisenfeld2001, Bottner2003, Dumke2004,Galliano2003, Galliano2005, Marleau2006, Bendo2006}. This excess can be interpreted as very cold dust ($\\le$10K), in which case more than 50$\\%$ of the total dust mass of these galaxies should reside in a very cold component. {\\revisedbis Cold dust is also needed to explain the break in the Gas to Dust mas ratio as a function of metallicity relaion for low-metallicity galaxies~\\citep{Galliano_Dwek_Chanial_2008, Munoz2009}}. The presence of this cold dust component is still a contentious issue in the ISM community and will have important consequences on our comprehension of ISM properties of low metallicity environments. \\citet{Lisenfeld2001}, \\citet{Reach1995}, \\citet{Dumke2004}, \\citet{Bendo2006} or \\citet{Meny2007} suggested that changes in dust emission properties (changes in dust emissivity or resonances related to dust impurities) should be responsible for boosting submm emission above the 15-20K thermal emission expected at these wavelengths. However, not all low metallicity galaxies show submm excess, as shown recently by the observations of the nearby Local Group Galaxy {\\revisedbis IC10 {\\revisedbis (Parkin et al. 2009 in prep)}, where the main} two star forming regions were isolated with ISO, \\spitz\\ and 850 \\mic\\ observations, the SEDs were modeled without invoking a very cold dust component. Moreover, \\citet{Draine2007} showed that their observations of mostly metal-rich galaxies can largely be reproduced by dust models which do not account for a very cold dust component, even in their limited number of cases where submm observations are present. Studies using submm observations for a wider range of metallicity values are necessary to check the relevance of these conclusions for low metallicity environments.  In this paper, we present the first \\APEX/ \\lab\\ 870 \\mic\\  observations of dwarf galaxies: 1 extended galaxy (NGC~1705) and 3 compact sources (Haro~11, UM~311, Mrk~1089) (see Table~\\ref{Sample_properties}). We have combined these data with \\spitz\\ and/or \\iras\\ observations to produce global SEDs that we use to model the dust properties. The sample is small but covers a wide range of metallicities, from $\\sim$1/9 \\zsun\\ for Haro~11 to $\\sim$1/3 \\zsun\\  for NGC~1705. It also presents varied morphologies, size scales and characteristics: resolved or compact galaxies, disturbed and even merging environments.  We describe the sample in $\\S$ 2 and the observations and data reduction in $\\S$ 3 and the images and photometry in $\\S$ 4. In $\\S$ 5, we present the SED modelling and discuss the results in $\\S$ 6. ",
        "conclusions": "The quantification of the dust mass of a galaxy aids our understanding of its evolution and star formation history. Larger dust masses are sometimes found in low metallicity galaxies when using submm constraints in the SED modelling. In this context, submm observations are clearly necessary to lead a more complete description of the distribution and properties of dust.  We focused our paper on the dust modelling of four low metallicity galaxies observed with \\APEX/\\lab. \\\\  In this paper:  \\begin{enumerate}  \\item We present the first images of four dwarf galaxies carried out with the \\APEX/\\lab\\ instrument observing at 870 \\mic.  \\item We construct the SEDs with \\spitz\\ IRAC and MIPS bands as well as the IRS spectra for Haro~11. We apply our SED model and determine  the dust properties of these galaxies.  \\item We find that the mass of PAHs accounts for 0.08 to 0.8 $\\%$ of the total dust mass of the galaxies, which is a factor of 5 to 50 lower than that of the Galaxy.  \\item To investigate the influence of the submm constraints on the interpretation, we test the effect on the SED model results when submm 870 \\mic\\ observations are taken into account and compare with SED models not taking into account the 870 \\mic\\ flux, but with observational constraints at wavelengths only as long as 160 \\mic. We find that the use of submm observational constraints always leads to an increase of  the total dust mass derived for our low metallicity galaxies.  \\item We choose to include an additional component to account for the excess submm emission of NGC~1705 and Haro~11. A cold dust component ($\\sim$ 10K) with a $\\beta$ emissivity index of 1 substantially improves the fit. We find at least 70$\\%$ of the total dust mass residing in a cold ($\\sim$ 10K) dust component for these two galaxies. We note that describing a cold component of $\\beta$=2 does not give very different $\\bar{\\chi}$$^2$ values, but would give unrealistically larger D/G. Our results however do not rule out the hypothesis of a change in dust emissivity as a function of wavelength proposed in recent studies \\citep[e.g.][]{Dupac2003,Meny2007}.  \\item While Haro~11 has a substantial ($\\sim$ 70$\\%$) cold dust component, it also harbours a significant fraction of dust mass (30$\\%$) in a warmer dust component ($>$ 25K). The SED peaks at unusually short wavelengts (36 \\mic), highlighting the importance of the warm dust.  \\item We determine the D/G for Mrk~1089, NGC~1705 and Haro~11 to be 1.9 $\\times$ 10$^{-3}$,  4.1 $\\times$ 10$^{-3}$ and 0.2, respectively. For Mrk~1089 and NGC~1705, these D/G are consistent with current chemical evolution models. On the contrary, Haro~11 has an excessively high D/G considering the upper limits detected in HI and CO. Haro~11 also falls far from the  Schmidt law, perhaps due to the observed deficit of gas in this galaxy. This could suggest the presence of a large amount of molecular gas. 10 times more molecular gas, compared to that deduced from CO measurements, may be present but not necessarily traced by CO observations.  \\item From our SED models, we determine the total infrared luminosity of our galaxy sample to range from 5.8 $\\times$ 10$^{7}$ for NGC~1705 to 1.7 $\\times$ 10$^{11}$ for the LIRG Haro~11. These values of L$_{TIR}$ are systematically higher than those obtained using the \\citet{Dale_Helou_2002} formula but compare better to the \\citet{Draine_Li_2007} formula. While $\\sim$90$\\%$ of the dust mass is residing in the FIR to submm regime, not more than 6$\\%$ of the total IR luminosity in Haro~11 emerges from the FIR to submm (100 to 1100 \\mic), while most of the luminosity (70$\\%$) emerges in the NIR to MIR (3 \\mic\\ to 50 \\mic) window. This is in contrast to Mrk~1089 and NGC~1705 which distribute their luminosities more equally in these two wavelength windows.   \\end{enumerate}  Better observational coverage of the Rayleigh Jeans side of the SEDs should help us to disentangle possible scenarios to explain the excess detected in some galaxies in the submm: dominant diffuse ISM dust, modifications of the dust optical properties at submm wavelengths, very cold dust component etc. The \\hersc\\  guaranteed time key program of 50 dwarf galaxies (PI: S.Madden), which covers a wide range of metallicity values, will provide a broader coverage of wavelengths between 60 \\mic\\ to 600 \\mic\\ and will enable us to better sample the warm and cold dust. The observations will especially enable us to study the slope of the Rayleigh Jeans side of the SED and to learn if an independant very cold dust component does indeed exist. A following paper will discuss the systematic increase in the dust mass estimate with or without submm constraints and the implications for the D/G ratios for a broader sample of galaxies and study how it could possibly be influenced by metallicity. \\\\"
    },
    "0910/0910.0004_arXiv.txt": {
        "abstract": "We use  the overdensity  field reconstructed  in the volume  of the  COSMOS area  to study  the nonlinear  biasing  of the zCOSMOS galaxies. The galaxy  overdensity field is reconstructed using the current sample of $\\sim$8500 accurate zCOSMOS redshifts at $I_{AB} <  22.5$ out  to z$\\sim$1  on scales  $R$  from 8  to 12  \\hh Mpc.  By comparing  the  probability  distribution  function  (PDF)  of  galaxy density contrast $\\delta_g$ to  the lognormal approximation of the PDF of  the mass  density contrast  $\\delta$, we  obtain the  mean biasing function $b(\\delta,  z,R)$ between  the galaxy and  matter overdensity field  and  its  second  moments  $\\hat{b}$  and $\\tilde{b}$ up  to  $z  \\sim  1$. Over  the redshift  interval $0.4<z<1$ the conditional mean function $\\langle \\delta_g|\\delta \\rangle  = b(\\delta,  z,R) \\delta$  is of the following characteristic shape. The function  vanishes in the most underdense  regions and then sharply  rises in  a  nonlinear  way towards  the  mean densities. $\\langle \\delta_g|\\delta \\rangle$ is almost a linear tracer of the matter in the overdense regions, up to  the most overdense regions  in which it is  nonlinear again and the local effective slope of $\\langle \\delta_g|\\delta \\rangle$ vs. $\\delta$  is smaller than unity. The $\\langle \\delta_g|\\delta \\rangle$ function is evolving  only  slightly over  the redshift  interval $0.4<z<1$. The  linear biasing  parameter increases from $\\hat{b}=1.24  \\pm 0.11$ at $z=0.4$ to  $\\hat{b}=1.64 \\pm 0.15$ at  $z=1$  for  the  $M_B<-20-z$  sample  of galaxies. $\\hat{b}$ does  not show  any  dependence on  the smoothing  scale from 8 to 12 \\hh Mpc,  but increases  with  luminosity.  The measured nonlinearity  parameter $\\tilde{b}/ \\hat{b}$ is  of the order of a  few percent (but it  can be consistent  with 0) and it  does not change  with  redshift, the smoothing  scale or  the luminosity. By matching the linear  bias of galaxies to the halo bias, we infer  that the $M_B<-20-z$  galaxies reside in dark  matter haloes with  a  characteristic mass  of  about  $3-6  \\times 10^{12}$  \\Msol, depending on the halo bias fit. ",
        "introduction": "The large scale structure in the Universe is believed to have formed via gravitational instability of small, primordial density fluctuations. Virialised dark matter haloes are produced by the collapse of some overdense regions followed by hierarchical merging. Galaxies are formed within the dark matter haloes through multiplex processes including gas cooling, star formation and feedback which are difficult to model accurately \\citep[e.g.][]{White&Rees.1978}. It is thus expected that the relation between galaxies and the underlying matter distribution will be also complex. In particular, the efficiency of galaxy formation and the rate of galaxy evolution might vary from place to place depending on the matter density field. Therefore the actual galaxy distributions is not expected to be a rightful tracer of the underlying mass. This phenomenon is often referred to as ``galaxy biasing''.  In order to extract cosmological information from galaxy surveys it is important to model and parameterise galaxy biasing. This often requires using a statistical approach. Assuming that galaxies preferentially form within the peaks of the primordial density distribution, \\citet{Kaiser.1984} showed that the two-point correlation function of the galaxy distribution, $\\xi_{gg}$, should be amplified with respect to the mass autocorrelation function, $\\xi_{mm}$, according to the relation: \\be \\xi_{gg}(r) = b_{\\xi}^2 \\xi_{mm}(r) \\label{eq_biasksi} \\ee where the ``biasing parameter'' $b_{\\xi}$ is independent of the spatial separation $r$. A similar relation is obtained by relating the density contrast of the matter $\\delta$ and of the galaxies $\\delta_g$ at some position $\\bf{r}$ through the deterministic and linear relation \\be  \\delta_g ({\\bf r}) = b\\, \\delta ({\\bf r}) ~. \\label{eq_biasdetlin} \\ee This is the simplest model for galaxy biasing and is still commonly used. While equation~\\ref{eq_biasksi} follows from equation~\\ref{eq_biasdetlin}, the opposite does not hold. An obvious deficiency in the definition of $\\delta_g$ above is that it will break down in the most underdense regions $\\delta\\ll0$ if $b>1$, as values of $\\delta_g<-1$ are not possible. This implies that the galaxy bias $b$ must be a nonlinear function of $\\delta$ and, in general, it can also vary with redshift $z$, galaxy type and the smoothing scale $R$ used to define the density contrast: \\be b=b(\\delta, z, R).  \\ee \\citet{Fry&Gaztanaga.1993} proposed to parameterise this function in terms of coefficients of the Taylor expansion \\be \\delta_g=b_0+b_1\\,\\delta+\\frac{b_2}{2}\\,\\delta^2+\\dots  \\;, \\label{eq_biasfg} \\ee which are not fully independent as the conditions $\\langle \\delta_g \\rangle=0$ and $\\delta_g(\\delta=-1)=-1$ must hold.   Galaxy biasing  is  also expected  to  have a stochastic element:  for  any  given  value of $\\delta$ there will be a whole distribution of values for $\\delta_g$. The stochasticity originates from a number of different sources. First, the dynamics of large scale flows depends on extra variables beyond the value of the local density contrast $\\delta$ (e.g. on the tidal tensor) and makes the bias relation nonlinear, non-local and stochastic \\citep{Catelan.etal.1998}. Second, the efficiency of galaxy formation depends on details of the gas physics. Third, galaxies are discrete objects and any attempt to reconstruct $\\delta_g$ will be effected by shot noise.  \\citet{Dekel&Lahav.1999} have  proposed a formalism  which  separately accounts for the nonlinearity  and stochasticity  of  the biasing process. Galaxy biasing is described in terms of the conditional probability function $P(\\delta_g|\\delta)$ and its moments. A key quantity here is the mean biasing function $b(\\delta)$ defined by the conditional mean: \\be b(\\delta) \\delta = \\langle \\delta_g | \\delta \\rangle = \\int d \\delta_g P(\\delta_g|\\delta) \\delta_g . \\ee The mean biasing function $b(\\delta$) and its nonlinearity can be characterised by its second non-trivial moments: \\be  \\hat{b} \\equiv \\frac {\\langle b(\\delta)\\delta^2 \\rangle}{\\sigma^2} \\ee and \\be \\tilde{b}^2 \\equiv \\frac {\\langle b^2(\\delta)\\delta^2 \\rangle}{\\sigma^2} , \\ee with  $\\sigma^2$ the variance of the mass density contrast distribution. The parameter $\\hat{b}$ measures the slope of the linear regression of $\\delta_g$ against $\\delta$. In the case of linear biasing (see equation~\\ref{eq_biasdetlin}), both $\\hat{b}$ and $\\tilde{b}$ reduce to the constant bias. The ratio $\\tilde{b}/\\hat{b}$ is thus a measure of the nonlinearity in the biasing relation. Moreover, the local variance of $\\delta_g$ at fixed $\\delta$, $\\sigma^2_g(\\delta)$ can be used to quantify the degree of stochasticity of the biasing relation.   Based on the Press-Schechter formalism and its extensions \\citep{Bond.etal.1991}, \\citet{Mo&White.1996} developed an analytical model for the mean biasing relation of the dark matter haloes. This assumes that large scale motions follow the spherical collapse approximation. The general case is discussed by \\citet{Catelan.etal.1998}. Related work has been presented in \\citet{Mo.etal.1997} and Porciani et al. (1998, see also Scannapieco \\& Barkana 2002) \\nocite{Porciani.etal.1998, Scannapieco&Barkana.2002} where two-point and higher-order statistics are considered. Following the analytical approach by \\citet{Mo&White.1996}, a number of studies based on N-body simulations were carried out to study the halo bias \\citep[e.g.][]{Jing.1998, Porciani.etal.1999,  Sheth&Lemson.1999, Sheth&Tormen.1999, Jing.1999, Kravtsov&Klypin.1999, Sheth.etal.2001, Seljak&Warren.2004, Tinker.etal.2005, Pillepich.etal.2008}, leading to a new set of fitting formulae for the mean biasing relation, and a better understanding of the origin of halo biasing. Independently of the exact halo definition, assumed cosmology, simulation box size and resolution, there is a consensus that in a cold-dark matter scenario: i) at a given epoch, more massive haloes are more biased tracers of the underlying matter than lower mass haloes; ii) for halos of fixed mass, the amount of biasing increases with redshift.  However, it is still a huge step from a successful description of ``halo biasing'' to that of ``galaxy biasing'' as the latter requires incorporating a recipe for galaxy formation (and evolution) within the current cosmological framework. Galaxy biasing has been studied through hydrodynamic simulations \\citep[e.g.][]{Blanton.etal.1999, Blanton.etal.2000, Cen&Ostriker.2000, Yoshikawa.etal.2001} and semi-analytical modelling combined with N-body simulations \\citep[e.g.][hereafter SBD]{Kauffmann.etal.1997, Benson.etal.2000, Somerville.etal.2001, Sigad.etal.2000}. Despite the difference in the treatment of the various gas-related processes,",
        "conclusions": "\\label{sec_concl}   In this work, we make use  of the reconstructed overdensity field in the zCOSMOS volume \\citep[see][]{Kovac.etal.2009} to derive the conditional mean function $\\langle \\delta_g|\\delta \\rangle = b(\\delta, z, R) \\delta$ and the second moments of the mean biasing function $b(\\delta, z, R)$. For this purpose we employ the density field on a grid reconstructed by using the three dimensional distances between galaxies and grid points and counting the objects within a spherical top-hat aperture. We implement a novel method ZADE \\citep{Kovac.etal.2009} to account for galaxies not yet observed spectroscopically in the selected samples of galaxies used to reconstruct the density field. For a biasing analysis, the main advantage of the ZADE method is that in a statistical sense, the mean intergalaxy separation is that of all galaxies in the selected galaxy sample, and not only of the sample of galaxies with spectroscopic redshifts.   We have carried out a number of tests on the mock catalogues to assess various errors which are going to affect our biasing analysis. Particularly, we have empirically estimated uncertainties due to cosmic variance, shot noise errors and the density field reconstruction errors. \\begin{itemize}  \\item Cosmic variance errors cause a spread in the $\\langle \\delta_g|\\delta \\rangle$ function: for a given $\\delta$ there is a range of $\\delta_g$ values measured in the mock catalogues. Quantifying this spread by the standard deviation $\\sigma$ of $\\log(1 + \\delta_g)$, we find that $\\sigma$ is largest in the most underdense regions where $\\sigma \\sim 0.1$, becomes lower at the intermediate $\\delta$ values, $\\sigma < 0.05$, and increaseas again in the most overdense regions up to $\\sigma \\sim 0.05$ at a given $\\delta$.   \\item The shot noise (discrete galaxy sampling) errors modify significantly the shape of the $\\langle \\delta_g|\\delta \\rangle$ function in the most underdense regions, making the local bias $b(\\delta, z, R)$ values in the same regions to appear higher.   \\item Reconstruction errors are relevant only in the underdense regions. At the reconstructed value of $\\log(1+\\delta_g)=-1$ in the galaxy density field, the reconstruction error can cause an uncertainty of the order of 0.1 in the matter density field $\\log(1+\\delta)$.  \\item The $\\hat{b}$ parameter increases due to the shot noise and reconstruction errors. The $\\tilde{b}/\\hat{b}$ parameter is not susceptible to either of these errors. The cosmic variance causes a spread in the measured values of both of these parameters.  \\end{itemize}   We can summarise our main findings in the biasing analysis of the 10k zCOSMOS galaxies as follows:  \\begin{itemize}  \\item The conditional mean function $\\langle \\delta_g|\\delta \\rangle$ has a characteristic shape as described below. In most underdense regions, the mean biasing function vanishes. At some $\\delta < 0$, the mean biasing function appears and then rises sharply in the underdense regions, with the local slope of the biasing function larger than unity. Starting from around mean density and towards higher overdensities, the $\\langle \\delta_g|\\delta \\rangle$ function closely follows a linear relation $\\delta_g = b \\delta$ with $b$ a constant. In the most overdense regions zCOSMOS galaxies are antibiased, i.e. locally $b(\\delta,z,R)<1$. This is true for all samples of tracer galaxies used. The conditional mean function $\\langle \\delta_g|\\delta \\rangle$ is clearly nonlinear in the most overdense and underdense regions.   \\item There is a detectable change in the shape of the $\\langle \\delta_g|\\delta \\rangle$ function with redshift in the overdense regions. For a given population,  galaxies become more biased tracers of the matter in the regions of mean and mildly positive overdensities as redshift increases from 0.4 to 1. There is an indication of an evolution in the value of $\\delta$ in the overdense regions, at which galaxies become antibiased. For a given population of tracer galaxies, this happens at higher $\\delta$ for higher $z$. Taking into account all the sources of errors, we cannot discriminate if the shape of the $\\langle \\delta_g|\\delta \\rangle$ function in the underdense regions stays constant or undergoes some evolution in $0.4<z<1$. It is worth stressing that redshift evolution may be not visible due to the overlapping redshift bins in the analysis.   \\item The $\\langle \\delta_g|\\delta \\rangle$ function shows some dependence on the scale in the most overdense regions, more pronounced at high redshift, such that at smaller scales galaxies are more biased tracers of the underlying matter distribution. However, this is not highly statistically significant.   \\item When comparing the $\\langle \\delta_g|\\delta \\rangle$ function for the different luminosity selected tracer galaxies, there is an indication that more luminous galaxies are more biased than less luminous galaxies in the regions from about mean density to the highly overdense regions.   \\item The linear bias between the 10k zCOSMOS galaxies and mass is increasing with redshift. There is some evidence that more luminous galaxies are more biased tracers of matter. We do not detect any significant dependence of the linear biasing parameter on the  scale at which we measure galaxy overdensity fields. The nonlinearity parameter is of the order of a few percent. It does not change with the redshift, with the smoothing scale or with the luminosity of the tracer galaxies.  \\item The linear bias of the stellar mass-weighted density field is either larger or the same within the errors compared to the linear bias of the $B$-band luminosity- or unity-weighted density field.   \\item By comparing the galaxy biasing to the halo biasing, using the approximation for halo bias of \\citet{Sheth.etal.2001} and \\citet{Pillepich.etal.2008}, we infer that the $M_B<-20-z$ zCOSMOS galaxies in $0.4<z<1$ reside in dark matter haloes with a characteristic mass of $\\sim 3$ or $\\sim 6 \\times  10^{12}$ \\hh \\Msol, using these two models, respectively. One would need to work with the stellar mass complete samples to obtain the evolution in characteristic halo mass whose physical interpretation is not ambiguos.    \\end{itemize}   Broadly speaking, our results are in line with findings from the previous study of the nonlinear biasing of high redshift galaxies \\citep{Marinoni.etal.2005, Marinoni.etal.2008} and, qualitatively, they follow the biasing history outlined by the theoretical works \\citep[e.g.][]{Blanton.etal.1999, Blanton.etal.2000, Somerville.etal.2001}. When going into details, there are a number  of discrepancies which neeed to be solved, such as the exact ``speed'' in the evolution of the linear biasing parameter or the dependence of biasing on the luminosity of tracer galaxies. While the current biasing results from the $z<1.5$ surveys are important as they support the general picture of the biased galaxy formation and provide a framework for the future theoretical work, the fine tuning of the galaxy formation models is still hampered by the limitations of the existing  spectroscopic surveys at these redshifts, particularly by the cosmic variance."
    },
    "0910/0910.0232_arXiv.txt": {
        "abstract": "We study high-energy gamma-ray afterglow emission from gamma-ray bursts (GRBs) in the prior emission model, which is proposed to explain the plateau phase of the X-ray afterglow. This model predicts the high-energy gamma-ray emission when the prompt GRB photons from the main flow are up-scattered by relativistic electrons accelerated at the external shock due to the prior flow. The expected spectrum has the peak of $\\sim 10$--100~GeV at around the end time of the plateau phase for typical GRBs, and high-energy gamma rays from nearby and/or energetic GRBs can be detected by current and future Cherenkov telescopes such as MAGIC, VERITAS, CTA, and possibly \\textit{Fermi}. Multi-wavelength observations by ground-based optical telescopes as well as \\textit{Fermi} and/or \\textit{Swift} sattelites are important to constrain the model. Such external inverse-Compton emission may even lead to GeV-TeV gamma-ray signals with the delay time of $\\sim 10$--100~s, only if the plateau phase is short-lived. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.2237_arXiv.txt": {
        "abstract": "We develop a variant of the generalized slow roll approach for calculating the curvature power spectrum that is well-suited for order unity deviations in power caused by sharp features in the inflaton potential.    As an example, we show that predictions for a step function potential, which has been proposed to explain order unity glitches in the CMB temperature power spectrum at multipoles $\\ell = 20-40$, are accurate at the percent level.   Our analysis shows that to good approximation there is a single source function that is responsible for observable features and that this function is simply related to the local slope and curvature of the inflaton potential.  These properties should make the generalized slow roll approximation useful for inflation-model independent studies of features, both large and small, in the observable power spectra. ",
        "introduction": "The ordinary slow roll approximation provides a model-independent technique for computing the initial curvature power spectrum for inflationary models where the scalar field potential is sufficiently flat and slowly varying.  Such models lead to curvature power spectra that are featureless and nearly scale invariant (e.g. \\cite{Lidsey:1995np}).  On the other hand, features in the inflaton potential  produce features in the power spectrum. Glitches in the observed temperature power spectrum of the cosmic microwave background (CMB) \\cite{Bennett:2003bz} have led to recent interest in exploring such models (e.g.~\\cite{Peiris:2003ff,Covi:2006ci,Hamann:2007pa,Mortonson:2009qv,Pahud:2008ae,Joy:2008qd}).    To explain the glitches as other than statistical flukes, these models require order unity variations in the curvature power spectrum across about an $e$-fold in wavenumber.  Such cases are typically handled by numerically solving the field equation on a case-by-case basis (e.g.~\\cite{Adams:2001vc}).   For model-independent constraints and model building purposes it is desirable to have a simple but accurate prescription that relates features in the inflaton potential to features in the power spectrum (cf.~\\cite{Hunt:2004vt,Habib:2004kc,Joy:2005ep,Kadota:2005hv}).  The generalized slow roll  (GSR) approximation was introduced by Stewart \\cite{Stewart:2001cd} to overcome some of the problems of the ordinary slow roll approximation for potentials with small but sharp features.  In this approximation, the ordinary slow roll parameters are taken to be small but not necessarily constant. In this paper we examine and extend the  GSR approach for the case of large features where the slow-roll parameters are also not necessarily small.  In \\S \\ref{sec:GSR}, we review the GSR approximation and develop the variant for large power spectrum features.   In the Appendix, we compare this variant to other GSR approximations in the literature \\cite{Stewart:2001cd,Choe:2004zg,Dodelson:2001sh,Kadota:2005hv,Gong:2005jr}.   We show that our variant provides both the most accurate results and is the most simply related to the inflaton potential. In \\S \\ref{sec:applications}, we show how this technique can be used to develop alternate inflationary models to explain a given observed feature.   We discuss these results in \\S \\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  We have shown that a variant of the generalized slow roll (GSR) approach remains percent level accurate at predicting  order unity deviations in the observable CMB temperature and polarization power spectra from sharp potential features.   Unlike other variants which explicitly require $|\\eta_{H}|\\ll 1$, and hence nearly constant $\\epsilon_{H}$, our approach allows $\\eta_{H}$ to be order unity, as long as it remains so for less than an $e$-fold, and hence $\\epsilon_{H}$ to vary significantly.  We have tested our GSR variant against a step function model that has been proposed to explain features in the CMB temperature power spectrum at $\\ell \\sim 20-40$.   Our analysis also shows that to good approximation a single function, $G'(\\ln \\eta)$, controls the observable features in the curvature power spectrum even in the presence of large features. We have explicitly checked this relationship and the robustness of our approximation by constructing two different inflationary models with similar $G'$.  Therefore observational constraints from the CMB can be mapped directly to constraints on this function independently of the model for inflation. Moreover, this function is also simply related to the slope and curvature of the inflaton potential in the same way that scalar tilt is related to the potential in ordinary slow roll $G' \\approx 3 (V_{,\\phi}/V)^{2} - 2 (V_{,\\phi\\phi}/V)$. These model independent constraints can then be simply interpreted in terms of the inflation potential.   We intend to explore these applications in a future work.  {\\it Acknowledgments:}  We thank Hiranya Peiris for sharing code used to crosscheck exact results.   We thank Michael Mortonson, Kendrick Smith and Bruce Winstein for useful conversations. This work was supported by the KICP under NSF contract PHY-0114422.   WH was additionally supported by DOE contract DE-FG02-90ER-40560 and the Packard Foundation.  \\appendix    \\begin{figure}[tbp] \\includegraphics[width=0.45\\textwidth]{ExactPowersp_ML_3ML.eps} \\caption{Curvature power spectrum for the ML and 3ML models.} \\label{plot:ExactPowersp_ML_3ML} \\end{figure}"
    },
    "0910/0910.5621_arXiv.txt": {
        "abstract": "{% We present projected rotational velocity measurements of the red giant in the symbiotic star MWC~560, using the high-resolution spectroscopic observations with the FEROS spectrograph. We find that the projected rotational velocity of the red giant is \\vsi $= 8.2 \\pm 1.5$~\\kms, and estimate its rotational period to be  $P_{\\rm rot}$ = 144 - 306~days. Using the theoretical predictions of tidal interaction and pseudosynchronization, we estimate the orbital eccentricity  $e=0.68-0.82$. We briefly discuss the connection of our results with the photometric variability of the object. } ",
        "introduction": "MWC~560 (V694~Mon) was discovered as an object with bright hydrogen lines (Merrill \\& Burwell 1943). It is a symbiotic binary system, which consists of a red giant and a white dwarf (Tomov et al. 1990; Michalitsianos et al. 1993). The orbital period is supposed to be $P_{\\rm orb} \\approx 5.3$~yr  (Doro-shenko, Goranskij \\& Efimov 1993). The most spectacular features of this object are the collimated ejections of matter with velocities of up to $\\sim 6000$~\\kms\\ (Tomov et al. 1992; Stute \\& Sahai 2009) and the resemblance of its emission line spectrum to that of the low-redshift quasars (Zamanov \\& Marziani 2002).  The jet ejections are along the line of sight and the system is seen almost pole-on ($i < 16^\\circ$). This makes it difficult to obtain the orbital parameters using the radial velocity variations, and impossible to observe effects such as eclipse or illumination.  To improve our understanding of this object, we aim here to: (1) measure the projected rotational velocity of the red giant (\\vsi) from high resolution spectra, and (2) use our measured \\vsi to find the eccentricity $e$ of the orbit, on the basis of the theory for pseudosynchronization in binary stars. ",
        "conclusions": "In most of the symbiotics the activity of the hot components is irregular/aperiodic. However, in MWC~560 it has a periodic character. An analysis of the historical light curve (Luthardt 1991) and B-band photometry, as well as the periodogram analysis of V band and near-IR observations reveal that this periodicity has remained in-phase for over a century (Doroshenko et al. 1993; Gromadzki et al. 2007). The most natural explanation for this phenomenon is orbital modulation: around periastron, the mass accretion rate increases, which causes an increase of the accretion luminosity and the brightness of the accretion disk, which produces the observed modulation of the light curves (see also Doroshenko et al. 1993; Gromadzki et al. 2007).   With the parameters assumed above, we can calculate the distance between the components at the periastron: $r_{\\rm p}=a(1-e) = 180 - 230~R_\\odot$. It follows that the red giant fills the Roche lobe for about  5\\% -- 25 \\% of the orbital period. This is in agreement with the suppositions of Gromadzki et al. (2007), that the orbital eccentricity and corresponding Roche lobe overflow at the periastron is the reason for the observed photometric variability.  It has been noted that the rotational period of the mass donor is considerably shorter than the orbital period (Zamanov et al. 2008) in a few jet-ejecting symbiotics. Our findings here pose the question whether their orbits are also eccentric."
    },
    "0910/0910.5598_arXiv.txt": {
        "abstract": "We present a pre-survey study of using Pan-STARRS  high sampling rate video mode guide star images to search for TNOs. Guide stars are primarily used by Pan-STARRS to compensate for image motion and hence improve the  point spread function. With suitable selection of the guide stars within the Pan-STARRS 7 $\\rm deg^{2}$ field of view, the lightcurves of these guide stars can also be used to search for occultations by TNOs. % The best target stars for this purpose are stars with high signal-to-noise ratio (SNR) and small angular size. In order to do this, we compiled a catalog using the SNR calculated  from stars with $m_{\\rm V} <13$  mag in the Tycho2 catalog then cross matched these stars with the 2MASS catalog and estimated their  angular sizes from $(V-K)$ color. We  also outlined a new detection method based on matched filter that is optimized to search for  diffraction patterns in the lightcurves due to occultation by sub-kilometer TNOs. A  detection threshold  is set to compromise between real detections and false positives. Depending on the theoretical size distribution model used, we expect to find up to \\emph{a hundred events} during the  three-year life time of the Pan-STARRS-1 project. The high sampling (30~Hz) of the project facilitates detections of small objects (as small as 400~m), which are numerous according to  power law size distribution,  and thus allows us to verify various models and further constrain our understanding of the structure in the outer reach of the Solar System. We have tested the detection algorithm and the pipeline on a set of  \\emph {engineering} data (taken at 10~Hz in stead of 30~Hz). No events were found within the engineering data, which is consistent with the small size of the data set and the theoretical models. Meanwhile, with a total of $\\sim 22$ star-hours video mode data ($|\\beta| < 10^{\\circ}$), we are able to set an upper limit of  $N(>0.5 \\rm ~km)\\sim 2.47\\times10^{10}$ deg$^{-2}$ at 95\\% confidence limit. ",
        "introduction": "\\label{sec:intro} Since the first Kuiper Belt object (not including Pluto)  was discovered in 1992  \\citep{jewitt1993}, more than 1000  trans-Neptunian objects  (TNOs) have been found\\footnote{See http://www.cfa.harvard.edu/iau/lists/TNOs.html for a list of these objects.}. A picture is emerging of the principal characteristics of the TNOs that is rich in information but leaves many unanswered questions. For example, a number of dynamical groups have been identified: classical, resonant,  scattered disk \\citep{Luu1997}, and more recently the extended scattered disk \\citep{gladman2002}. What  is the number and size distribution of the smaller objects of these dynamical groups? Is there an extension of the Kuiper Belt beyond 50 AU comprising bodies too small to have been detected in direct surveys?  Meanwhile, could a possible close stellar encounter in the early history of the Solar System \\citep{Allen2001, Trujillo2001, bernstein2004, Fuentes2008} be responsible for the mass deficit and the depletion of larger objects (D $>$ 150~km) beyond 50~AU?  The size distribution of these small TNOs provides important clues to the dynamical evolution of the early Solar System. Kilometer-size TNOs with 4\\% albedo are expected to have $M_{\\rm R} \\gtrsim30$, which is still way below the detection limit of the largest ground-based telescopes. Yet the vast majority of the TNO population is beyond the limit of direct observation --for example, the  Keck pencil beam survey \\citep{Chiang1999} has a detection limit  of $M_{\\rm R} \\approx 27.5$.  A number of authors \\citep{Bailey1976, Roques2000} have suggested the possibility of detecting TNOs by stellar occultation. An occultation manifests  itself as the shadow created by a TNO occulting a background star sweeping across the Earth, causing flux reduction in the  lightcurve. Hence, as opposed to direct observation, stellar occultation by TNOs provides a unique way of detecting kilometer to sub-kilometer size objects in the foreseeable future. Recently a few groups have attempted to search for occultations by TNOs \\citep{Roques2006,bickerton2008,Lehner2009,kiwi2008,Bianco2009,bickerton2009}, but no real detection has been confirmed yet.   Projects committed to search for TNOs by occultation either do not have adequate time resolution, observe too few stars with good SNR or do not have enough observing time. These projects do a blind search for occultations, and these events are rare due to the reasons mentioned above.  Therefore, the key to a successful detection is to have  high sampling ($\\ge$~20~Hz), high signal to noise ratio of background stars  (SNR $\\ge$ 80),  small stellar angular sizes ($\\leq$ 0.1~mas) and, most importantly, many star-hours of observations ($>$ a few hundred thousands).  The Panoramic Survey Telescope and Rapid Response System (Pan-STARRS) is a project consisting of 4 telescopes that  can cover over  6000 $\\rm deg^{2}$ per night or scan the whole visible sky from Hawaii   ($3\\pi$) in a week to a detecting limit $\\sim 24^{\\rm th}$ magnitude. Pan-STARRS prototype telescope (PS1) will monitor up to 60 guide stars in a very high sampling rate video mode. We have compiled a list which allows us to select guide stars from anywhere in the field; guide stars could be in any of the 64$\\times$64 OTCCDs (see $\\S$\\ref{sec:ps1}) and read at $\\sim30$~Hz. Each field will have multiple choices of guide stars that were ranked based on the predicted event rates, which depend on their SNRs and angular sizes. By properly selecting guide stars,  the Pan-STARRS video mode images  would be ideal for searching TNOs.  In the next section, we briefly describe the Pan-STARRS system. In $\\S$\\ref{sec:diff} and $\\S$\\ref{sec:det} we discuss the diffraction profiles and present our detection algorithm. In $\\S$ \\ref{sec:angular} we describe the prediction of stellar angular size using $(V-K)$ color. In $\\S$\\ref{sec:rate} we estimate the number of expected events based on different number density estimations, sampling rates and stellar angular sizes.  The methodology and  compilation of  the guide star list is in  $\\S$\\ref{sec:catalog}.  In $\\S$\\ref{sec:engdata}, we show and describe the quality of the engineering data obtained in fall 2008. Conclusions are in the final section. ",
        "conclusions": "\\label{sec:result} We have made a pre-survey study of using lightcurves from guide star video mode images to search for occultation by TNOs near 43~AU. Simulations were made to calculate the null hypothesis distribution (\\fig{fig:nulldist1}) and  detection efficiency for various angular sizes and SNRs. Under 30~Hz sampling, PS1 can detect objects as small as $\\sim400$ m at 43~AU (\\fig{fig:efficiency}). Using $(V-K)$ indices, we predicted the angular sizes of all matched stars between Tycho2 and 2MASS catalogs with $m_{\\rm V}<13$ mag and $m_{\\rm K}<16$ mag above the PS1 southern declination limit $\\delta \\ge-30^{\\circ}$. The distribution of stellar  angular sizes peak around 0.02~mas in PS1 $3\\pi$ sky (\\fig{fig:vmk}). On average, there are about 180 stars with $m_{\\rm V}<11.5$ mag for PS1 7 deg$^{2}$ field of view.  Given that we can choose only 60 guide stars for each PS1 target field,  we ranked our guide star candidates  by their number of expected events based on their angular sizes, SNRs and model by P\\&S05. As mentioned before, multiple choices of guide star will ensure us the freedom of selecting guide stars away from gaps in the CCD array or dead pixels that would develop over time. Based on the differential size distribution from different models and the threshold  to have one false positive in the PS1 three-year lifetime, we estimated the number of expected events could be somewhere from 1 to $\\sim100$ (\\fig{fig:rate1}).  The engineering data allowed us to investigate the quality of the lightcurve and develop the detection pipeline for the upcoming real data. We have established that the detection technique performs as well with the filtered engineering data. We also realize that the true SNR is limited by systematics. Some of the systematics will be removed as we move to real data stream, however some of the systematics will remain. For example, scintillation will limit the SNR to about 200. We have recalculated the event rates with the worst case scenario and  found that the event rates were compromised by a factor of four. Even with this pessimistic estimation the event rate that PS1 will find can allow us to place a constraint on the size distribution and hence the evolution of the TNOs. Using the available engineering data and detection efficiency at 0.4 and 0.5~km, we were able to derive an effective solid angle $\\sim 2.3\\times 10^{-11}$ and $\\sim 1.2\\times 10^{-10}$ deg$^{2}$ and set the 95\\% confidence upper limit on surface number density at $N(>0.4 \\, \\rm km)\\sim 1.3\\times10^{11}$  and $N(>0.5 \\, \\rm km)\\sim 2.47\\times10^{10}$ deg$^{-2}$  (\\fig{fig:upper}). In future work, for the upcoming real data, we will set the threshold based on  maximum true positive to false positive ratio, which allows more candidate events for further investigations. Meanwhile, the video mode lightcurves can also be used to search for objects in Sedna like orbits from 100 to 1000~AU. We will work on a new detection algorithm that is capable of searching for objects in this region."
    },
    "0910/0910.2976_arXiv.txt": {
        "abstract": "{ Since a century cosmic rays are based on direct cosmic particle detection in space (below PeV)  or on secondary downward vertical airshowers (above TeVs). We consider  the guaranteed physics of horizontal (hadron) air-showers, HAS,  developing at high ($30-40$) km altitudes, above and below these energy windows. Their morphology and information traces are different from vertical ones. Hundreds of km long HAS are often split by geomagnetic fields in a long (fan-like) showering with a twin spiral tail.  The horizontal fan-like airshowers are really  tangent and horizontal only at North and South poles. At different latitude their showering plane are turned and inclined by geo-magnetic fields. In particular at magnetic equator such tangent horizontal East-West airshowers are bent and developed into a vertical fan air-shower, easily  detectable by a vertical array detector (hanging elements by gravity). Such \\emph{medusa } arrays maybe  composed  by inflated floating balloons chains. The light gas float and it acts as an calorimeter for the particles, while it partially  sustains the detector weight.  Vertically hanging chains  as well as  rubber doughnut balloons ( whose  interior may record Cherenkov lights)  reveal bundles of crossing electron pairs.  Even in space orbit such vertical array may hang by tiny tidal forces within huge balloon arrays, while brought in locus  by an extendable measure tape. Possibly located around Space Station in synergy with future AMS-particle detector.  Such an array  maybe loaded at best and cheapest prototype  in common balloons tracing upward and tangent hadron air-showers from terrestrial atmosphere edge. These array structure may reveal the split shower signature. Offering a way to disentangle better their shower origination, energy and interaction point. Better revealing the composition nature. Just beyond the Earth horizons there are exciting, but rarer, sources of upward airshowers: the new UHE Tau Air-showers astronomy originated within Earth by EeVs  tau  neutrino signals skimming the soil and forming UHE Tau, decaying later in flight. Therefore Horizontal airshowers at equator may show the hadron horizontal twin split nature, its composition as well as a very first expected  UHE Neutrino astronomy.\\ }  \\FullConference{European Physical Society Europhysics Conference on High Energy Physics, EPS-HEP 2009,\\\\ July 16 - 22 2009\\\\ Krakow, Poland}  \\begin{document} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.0468_arXiv.txt": {
        "abstract": "Giant planet formation process is still not completely understood. The current most accepted paradigm,  the core instability model, explains several observed properties of the solar system's giant planets but, to date, has faced difficulties to account for a formation time shorter than the observational estimates of protoplanetary disks' lifetimes, especially for the cases of Uranus and Neptune.  In the context of this model, and considering a recently proposed primordial solar system orbital structure, we performed numerical calculations of giant planet formation.  Our results show that if accreted planetesimals follow a size distribution in which most of the mass lies in 30-100 meter sized bodies, Jupiter, Saturn, Uranus and Neptune may have formed according to the nucleated instability scenario.  The formation of each planet occurs within the time  constraints and they end up with  core masses in good agreement with present estimations. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.4282_arXiv.txt": {
        "abstract": "We present a coherent and homogeneous multi-line study of the CS molecule in nearby (D$<$10Mpc) galaxies. We include, from the literature, all the available observations from the $J=1-0$ to the $J=7-6$ transitions towards NGC~253, NGC~1068, IC~342, Henize~2-10, M~82, the Antennae Galaxies and M~83. We have, for the first time, detected the CS(7-6) line in NGC~253, M~82 (both in the North-East and South-West molecular lobes), NGC 4038, M~83 and tentatively in NGC~1068, IC~342 and Henize~2-10. We use the CS molecule as a tracer of the densest gas component of the ISM in extragalactic star-forming regions, following previous theoretical and observational studies by Bayet et al. (2008a,b and 2009). In this first paper out of a series, we analyze the CS data sample under both Local Thermodynamical Equilibrium (LTE) and non-LTE (Large Velocity Gradient-LVG) approximations. We show that except for M~83 and Overlap (a shifted gas-rich position from the nucleus NGC~4039 in the Antennae Galaxies), the observations in NGC~253, IC~342, M~82-NE, M~82-SW and NGC~4038 are not well reproduced by a single set of gas component properties and that, at least, two gas components are required. For each gas component, we provide estimates of the corresponding kinetic temperature, total CS column density and gas density. ",
        "introduction": "\\label{sec:intro}   Detection of star-forming gas is one of the most direct ways to measure the star formation rate and activity in a galaxy, allowing us to significantly improve our understanding of galaxy formation and evolution. Star-forming regions of very dense gas (n(H$_{2}$)$> 10^{5}$cm$^{-3}$) are needed to maintain star formation activity, even in hostile environments associated with young massive stars. In these environments, star-forming regions are able to resist the disruptive forces (winds or radiation) from nearby newly formed stars longer than the local gas in the local interstellar medium \\citep{Klei83, LaRo83}. Determining the physical conditions of the very dense gas in which massive stars form is, thus, essential. In this paper, we study such dense gas and estimate its properties over a large range of nearby galaxy types.  Following theoretical studies by \\citet{Baye08a} and the first detections of extragalactic very dense gas presented in \\citet{Maue89a, Maue89b, Walk90, Mart05, Mart06b, Mart06a, Baye08b, Grev09}, we have carried out multi-line observations of the CS molecule in nearby (D$<$10Mpc) extragalactic environments, enhancing significatively the current data set of extragalactic CS observations. Indeed, CS lines have not been observed so far but in few brightest nearby nuclei such as NGC~253, IC~342 and M~82 and mainly in their lower-J rotational levels (see references above). Hence, the physical properties of such galaxies namely the kinetic temperature, gas density, etc were estimated using only a small sample of CS lines. Gas traced by higher-J of species such as CS has not be characterized so far. In addition, even for the brightest sources (e.g. NGC~253 and M~82), no study of the various velocity components nor various positions in the same galaxy has been performed so far in a systematic way. In this paper, first of a series, we thus aim at investigating in much more details the very dense gas properties in many extragalactic sources as traced by the CS molecule.  Sulphur-bearing species are shown to be particularly enhanced during massive star formation, while species such as HCN, although a useful dense gas tracer, may not be tracing the sites where star formation occurs \\citep{Lint06}. Recently, \\cite{Mart05} showed that sulfur emission in the nuclear region of the nearby starburst galaxy NGC~253 is very strong. Amongst sulfur-bearing species, the CS molecule appears as one of the best tracers of dense gas, especially its high-J rotational transitions such as the CS $J=5-4$ line with an excitation threshold higher than 10$^{4}$-10$^{5}$ cm$^{-3}$ \\citep{Bron96} and the CS $J=7-6$ line with a critical density of n$_{crit} \\sim 2\\times 10^{7}$ cm$^{-3}$ \\citep{Plum92}.  Selection of nearby sources where we have proposed to observe the CS lines was made through a comparison with past and recent CS observations (see references above). Sources where no CS detections were available from the literature (e.g. Henize 2-10) were chosen on the basis of the detections of the CO(6-5) or CO(7-6) lines since these lines are also good tracers of a relatively dense and quite warm gas \\citep{Baye04, Baye06}. The source selection also aimed at presenting various nearby galaxy types where CS line emission can be studied and compared. This sample is far from being unbiased but it gathers so far the most complete CS data sample ever obtained in extragalactic sources.  Thus, the centers of NGC~253, IC~342, Henize~2-10, the two molecular lobes of M~82 (North-East and South-West positions), three positions in the Antennae Galaxies (the two nuclei: NGC~4038 and NGC~4039, and a shifted position from NGC~4039 called Overlap) and the center of M~83 were selected, covering a large range of galaxy types and star-formation activities (starburst, irregular, merging galaxies, more quiet star-forming galaxies, etc). We have also included in our source sample the center of NGC~1068 (AGN-dominated galaxy) where HCN has been detected \\citep{Plan91,Krip08}.  We present our observations, data reduction and results in Sect. \\ref{sec:obs}. In Sect. \\ref{sec:mol} we analyze the data set both under Local Thermodynamical Equilibrium (LTE) and non-LTE (Large Velocity Gradient-LVG) approximations (see Subsects. \\ref{subsec:lte} and \\ref{subsec:lvg}). In Sects. \\ref{sec:disc} and \\ref{sec:con} we discuss our findings and conclude, respectively. ",
        "conclusions": "\\label{sec:con}  In order to better determine the properties of the very dense star-forming gas over a large range of physical conditions, we present the most complete, extragalactic CS survey ever performed on nearby galaxies. In particular, we detect for the first time the CS(5-4) and CS(7-6) lines, in various environments such as starburst, AGN-dominated galaxies, irregular galaxies, merging galaxies, etc.  In this first paper out of a series, we simply analyzed the data through a rotational diagram method and an LVG modelling and we show that for most sources at least two gas components are needed to reproduce the observations. The low-temperature gas component properties, derived from the CS(2-1), CS(3-2) and CS(4-3) lines (low-J CS lines), vary from source to source within the range T$_{K}$, n(H$_{2}$), N(CS) = 10-30K, 0.16-1.60$\\times 10^{5}$cm$^{-3}$, 0.25-6.30$\\times 10^{14}$cm$^{-2}$, respectively. The high-temperature gas component properties, derived from the CS(5-4) and CS(7-6) lines (high-J CS lines), vary from source to source within the range T$_{K}$, n(H$_{2}$), N(CS) = 45-70K, 6.30-40.0$\\times 10^{5}$cm$^{-3}$, 0.16-1.00$\\times 10^{14}$cm$^{-2}$, respectively. We also show that, using the current data set of CS observations, in Overlap and M~83, it appears that the very dense gas may be more homogeneously distributed than in other nearby sources.However, we again underline here that our estimates of the physical properties of these galaxies suffer from LVG degeneracy.  Comparison with other tracers of dense gas such as HCN, HNC, HCO$^{+}$, etc even in the brightest sources of our sample is very difficult to perform. However, in the molecular lobes of M~82, we have been able to show a good correlation between the CS and the methanol. In both NGC~253-1 and NGC~253-2, we showed that the HCN and the low-temperature CS gas on one side, and the HCO$^{+}$, H$_{2}$CO, HNC and the high-temperature CS gas on the other side have compatible rotational temperatures and column densities.  In pure AGN-dominated galaxy such as NGC~1068, we find no correlation between the HCN and CS.  Forthcoming papers will focuss on detailed studies of the very dense gas properties in individual nearby sources, investigating for instance their dense gas star-formation efficiencies, their dense gas masses based on CS modelling, etc, as well as on a more statistical analysis such as Kennicutt-Schmidt laws in both extragalactic and galactic sources.   \\begin{landscape} \\begin{table} \\caption{Observational parameters and Gaussian fits for the data set of CS observations. Where data were not available, we have put a black dash in the corresponding place in the table.}\\label{tab:obs} \\resizebox{18.8cm}{!}{ \\begin{tabular}{l c c c c c c c c c c c c c} \\hline Source & Line & $\\nu$ & Tsys & beam & Telesc. & $\\int$(T$_{mb}$ dv) & V$_{peak}$ & FWHM & T$_{peak}$ & rms & RA(J2000) & DEC(J2000) & Refs$^{1}$\\\\ & & & & size & name & & & & & & & & \\\\ & & (GHz) & (K) & ($''$) &  &  (Kkms$^{-1}$) & (kms$^{-1}$) & (kms$^{-1}$) & (mK) & (mK) & (h:m:s) & ($^{\\circ}$ : ' : '') &\\\\ \\hline NGC 253-1  & CS(2-1) & 97.980 & - & 51.0 & SEST & 4.4$\\pm$0.5 & 172 & 96 & 43.4 & 5.2 & 00:47:33.4 & -25:17:23.0 & d\\\\ & CS(3-2) & 146.969 & - & 16.7 & IRAM-30m & 11.9$\\pm$0.2 & 185 & 100 & 111.2 & 4.1 & '' & '' & d\\\\ & CS(4-3) & 195.954 & - & 12.6 & IRAM-30m & 11.5$\\pm$0.4 & 185 & 100 & 108.2 & 7.0 & '' & '' & d\\\\ & CS(5-4) & 244.936 & - & 21.0 & SEST & 4.4$\\pm$0.4 & 158 & 107 & 38.2 & 10.8 & '' & '' & d\\\\ & CS(7-6) & 342.883 & 679 & 14.0 & JCMT  & 7.9$\\pm$0.3 & 185$^{2}$ & 110.1$\\pm$4.4 & 67.1 & 5.1 & '' & '' & a\\\\ \\hline NGC 253-2  & CS(2-1) & 97.980 & - & 51.0 & SEST & 8.2$\\pm$0.5 & 290 & 116 & 66.5 & 5.2 & 00:47:33.4 & -25:17:23.0 & d\\\\ & CS(3-2) & 146.969 & - & 16.7 & IRAM-30m & 13.7$\\pm$0.2 & 288 & 117 & 110.3 & 4.1 & '' & '' & d\\\\ & CS(4-3) & 195.954 & - & 12.6 & IRAM-30m & 13.4$\\pm$0.5 & 288 & 121 & 104.5 & 7.0 & '' & '' & d\\\\ & CS(5-4) & 244.936 & - & 21.0 & SEST  & 5.4$\\pm$1.3 & 262 & 108 & 46.7 & 10.8 & '' & '' & d\\\\ & CS(7-6) & 342.883 & 679 & 14.0 & JCMT  & 4.4$\\pm$0.3 & 288$^{2}$ & 107.1$\\pm$7.1 & 38.7 & 5.1 & '' & '' & a\\\\ \\hline NGC 1068  & CS(3-2) & 146.969 & 600 & 16.7 & IRAM-30m & 9.1$\\pm$1.5 & 1100 & 245.0 & 30.0 & 10.0 & 02:42:40.7 & -00:00:47.6 & e\\\\ & CS(5-4) & 244.936 & - & 10.0 & IRAM-30m & 3.3$\\pm$0.3 & 1125$\\pm$9.0 & 180.0$\\pm$20.0 & 17.6 & 3.0 & '' & '' & f\\\\ & CS(7-6) & 342.883 & 307 & 14.0 & JCMT  & 1.4$\\pm$0.5 & 1125.4$\\pm$49.0 & 279.0$\\pm$99.3 & 4.8 & 8.1 & '' & '' & a\\\\ \\hline IC 342    & CS(1-0) & 48.991 & 350-550 & 36.0 & NRO & 2.5$\\pm$0.4 & 30 & 50-60 & 50.0 & 12.0 & 03:46:48.3 & 68:05:46.0 & g\\\\ & CS(2-1) & 97.980 & 188 & 25.1 & IRAM-30m & 5.0$\\pm$0.1 & 31.7$\\pm$0.6 & 53.9$\\pm$1.5 & 87.4 & 2.5 & '' & '' & c\\\\ & CS(3-2) & 146.969 & 225 & 16.7 & IRAM-30m & 4.9$\\pm$0.1 & 31.2$\\pm$0.6 & 51.7$\\pm$1.6 & 88.6 & 3.0 & '' & '' & c\\\\ & CS(5-4) & 244.936 & 378 & 10.0 & IRAM-30m  & 2.5$\\pm$0.2 & 38.2$\\pm$2.0 & 63.4$\\pm$5.1 & 37.6 & 2.9 & '' & '' & c\\\\ & CS(7-6) & 342.883 & 525 & 14.0 & JCMT  & 1.0$\\pm$0.4 & 29.3$\\pm$15.3 & 75.6$\\pm$30.6 & 12.3 & 8.5 & '' & '' & a\\\\ \\hline Henize 2-10 & CS(7-6) & 342.883 & 252 & 14.0 & JCMT  & 0.6$\\pm$0.2 & 883.1$\\pm$7.4 & 42.0$\\pm$16.1 & 13.0 & 7.4 & 08:36:15.2 & -26:24:34.0 & a\\\\ \\hline M~82\\_NE  & CS(2-1) & 97.980 & 236 & 25.1 & IRAM-30m & 9.4$\\pm$0.2 & 299.9$\\pm$1.3 & 103.4$\\pm$3.3 & 85.7 & 8.5 & 09:55:54.4 & 69:40:54.6 & a\\\\ & CS(3-2) & 146.969 & 278 & 16.7 & IRAM-30m & 8.9$\\pm$0.1 & 290.9$\\pm$0.7 & 114.6$\\pm$2.0 & 73.3 & 1.7 & '' & '' & a\\\\ & CS(4-3) & 195.954 & 1425 & 12.6 & IRAM-30m & 6.5$\\pm$0.8 & 309.2$\\pm$6.1 & 87.7$\\pm$12.1 & 70.1 & 17.2 & '' & '' & a\\\\ & CS(5-4) & 244.936 & 772 & 10.0 & IRAM-30m  & 6.4$\\pm$0.2 & 313.0$\\pm$1.1 & 73.5$\\pm$2.6 & 82.1 & 4.8 & '' & '' & a\\\\ & CS(7-6) & 342.883 & 309 & 14.0 & JCMT  & 1.1$\\pm$0.2 & 313.2$\\pm$4.4 & 58.1$\\pm$6.8 & 15.7 & 2.9 & '' & '' & a\\\\ \\hline M~82\\_SW & CS(2-1) & 97.980 & 237 & 25.1 & IRAM-30m & 8.9$\\pm$0.2 & 131.4$\\pm$1.2 & 109.8$\\pm$2.7 & 75.9 & 3.8 & 09:55:49.4 & 69:40:39.6 & a\\\\ & CS(4-3) & 195.954 & 1628 & 12.6 & IRAM-30m & 7.3$\\pm$1.0 & 140.0$\\pm$6.6 & 104.7$\\pm$17.2 & 65.5 & 17.9 & '' & '' & a\\\\ & CS(7-6) & 342.883 & 340 & 14.0 & JCMT & 1.0$\\pm$0.2 & 110.2$\\pm$8.8 & 60.3$\\pm$14.8 & 13.8 & 5.1 & '' & '' & a\\\\ \\hline NGC 4038  & CS(2-1) & 97.980 & 201 & 25.1 & IRAM-30m & 1.0$\\pm$0.1 & 1646.8$\\pm$3.0 & 79.6$\\pm$6.7 & 11.6 & 2.2 &  12:01:52.8 & -18:52:05.3 & a\\\\ & CS(3-2) & 146.969 & 471 & 16.7 & IRAM-30m & 0.8$\\pm$0.1 & 1645.2$\\pm$3.8 & 62.2$\\pm$7.6 & 11.4 & 2.6 & '' & '' & a\\\\ & CS(4-3) & 195.954 & 1396 & 12.6 & IRAM-30m & 1.1$\\pm$0.3 & 1646.0$\\pm$8.6 & 65.4$\\pm$15.8 & 15.3 & 6.5 & '' & '' & a\\\\ & CS(5-4) & 244.936 & 326 & 20.0 & JCMT  & 1.6$\\pm$0.2 & 1649.9$\\pm$6.5 & 103.2$\\pm$12.7 & 14.4 & 6.4  & '' & '' & b\\\\ & CS(5-4) & 244.936 & 1924 & 10.5 & IRAM-30m  & 1.9$\\pm$0.4 & 1666.3$\\pm$9.7 & 103.0 & 17.1 & 9.5 & '' & '' & a\\\\ & CS(7-6) & 342.883 & 250 & 14.0 & JCMT  & 0.7$\\pm$0.1 & 1629.4$\\pm$7.6 & 88.6$\\pm$15.0 & 7.8 & 2.7 & '' & '' & a\\\\ \\hline NGC 4039  & CS(2-1) & 97.980 & 226 & 25.1 & IRAM-30m & 1.7$\\pm$0.2 & 1640.8$\\pm$16.3 & 228.2$\\pm$33.6 & 7.1 & 3.2 & 12:01:53.5 & -18:53:11.3 & a\\\\ & CS(3-2) & 146.969 & 547 & 16.7 & IRAM-30m & 1.3$\\pm$0.3 & 1691.6$\\pm$19.1 & 145.0$\\pm$30.2 & 8.3 & 5.8 & '' & '' & a\\\\ & CS(4-3) & 195.954 & 1275 & 12.6 & IRAM-30m & - & - & - & - & 17.9 & '' & '' & a\\\\ & CS(5-4) & 244.936 & 277 & 20.0 & JCMT  & 1.4$\\pm$0.2 & 1757.6$\\pm$12.7 & 157.1$\\pm$26.5 & 8.2 & 4.0 & '' & '' & a\\\\ \\hline OVERLAP   & CS(2-1) & 97.980 & 203 & 25.1 & IRAM-30m & 0.9$\\pm$0.2 & 1498.8$\\pm$5.6 & 81.2$\\pm$18.9 & 10.6 & 2.7 & 12:01:54.9 & -18:52:59.0 & a\\\\ & CS(3-2) & 146.969 & 467 & 16.7 & IRAM-30m & 0.8$\\pm$0.2 & 1504.8$\\pm$6.4 & 61.8$\\pm$18.7 & 11.7 & 4.8 & '' & '' & a\\\\ & CS(4-3) & 195.954 & 1290 & 12.6 & IRAM-30m & - & - & - & - & 14.0 & '' & '' & a\\\\ & CS(5-4) & 244.936 & 1763 & 10.0 & IRAM-30m & - & - & - & - & 14.4 & '' & '' & a\\\\ & CS(7-6) & 342.883 & 230 & 14.0 & JCMT  & - & - & - & - & 4.3 & '' & '' & a\\\\ \\hline M~83      & CS(3-2) & 146.969 & 600 & 16.7 & IRAM-30m & 1.2$\\pm$0.3 & 553.0$\\pm$11.0 & 78.0$\\pm$24.0 & 15.0 & - & 13:36:59.2 & -29:52:04.5 & e\\\\ & CS(5-4) & 244.936 & - & 10.0 & IRAM-30m  & 2.3$\\pm$0.3 & 504.0$\\pm$7.0 & 82.4$\\pm$1.1 & 26.5 & 0.5 & '' & '' & d\\\\ & CS(7-6) & 342.883 & 121 & 14.0 & JCMT  & 0.4$\\pm$0.1 & 540.0$\\pm$9.5 & 99.1$\\pm$18.8 & 3.3 & 1.1 & '' & '' & a\\\\ \\hline \\end{tabular}}  $^{1}$ Refs: a: This work, b: See \\citet{Baye08b}, c: See \\citet{Alad09}, d: \\citet{Mart05}, e: \\citet{Maue89b}, f: \\citet{Mart09} and g: \\citet{Pagl95}; $^{2}$: The velocity positions of the two components in NGC 253 have been fixed to 185 kms$^{-1}$ and 288 kms$^{-1}$, consistently to what it has been performed in \\citet{Mart05}. \\end{table} \\end{landscape}   \\begin{table*} \\caption{Results of the linear regressions (slope, correlation coefficient, total source-averaged CS column density and rotational temperature) for a single- and two-components fit. The linear regressions have been obtained using the xmgrace software, including in the calculations the error bars of the observations.}\\label{tab:T_rot} \\begin{tabular}{ccccc} \\hline Source (Nb. of & slope & Correl. & N(CS)$^{1}$ &  T$_{rot}$ $^{1}$\\\\ fit components) & ($\\times 10^{-2}$) & Coeff. & ($\\times 10^{14}$) & (in K)\\\\ \\hline NGC 253-1 (one) & -2.67 & 0.878 & 1.79 & 16.3\\\\ NGC 253-1 (two) & -7.14 & 0.970 & 2.65 & 6.1 \\\\ & -1.33 & 0.969 & 0.80 & 32.7 \\\\ \\hline NGC 253-2 (one) & -3.40 & 0.910 & 2.47 & 12.8\\\\ NGC 253-2 (two) & -8.34 & 0.952 & 4.39 & 5.2\\\\ & -2.14 & 0.999 & 0.97 & 20.3 \\\\ \\hline NGC 1068 (one)  & -3.13 & 0.921 & 2.88 & 13.9\\\\ NGC 1068 (two)  & -6.09 & 1 & 5.80 & 7.1\\\\ & -1.31 & 1 & 0.68 & 33.2\\\\ \\hline IC 342 (one) & -4.26 & 0.936 & 2.67 & 10.2\\\\ IC 342 (two) & -10.87 & 0.997 & 2.87 & 4.0 \\\\ & -1.50 & 1 & 0.27 & 28.9 \\\\ \\hline M82-NE (one) & -3.91 & 0.958 & 2.64 & 11.1\\\\ M82-NE (two) & -7.81 & 0.992 & 4.25 & 5.6\\\\ & -2.77 & 0.999 & 1.04 & 15.7\\\\ \\hline M82-SW (one) & -3.94 & 0.975 & 1.80 & 11.0\\\\ M82-SW (two) & -7.10 & 1 & 2.39 & 6.1\\\\ & -3.02 & 1 & 0.74 & 14.4\\\\ \\hline size=7 $''$ & & & & \\\\ NGC 4038 (one) & -2.38 & 0.895 & 0.25 & 18.3 \\\\ NGC 4038 (two) & -6.24 & 0.936 & 0.29 & 7.0\\\\ & -1.52 & 0.991 & 0.15 & 28.6\\\\ size=25\\%$\\times$7 $''$ & & & & \\\\ NGC 4038 (one) & -2.43 & 0.891 & 0.42 & 17.9 \\\\ NGC 4038 (two) & -6.49 & 0.943 & 0.51 & 6.7\\\\ & -1.48 & 0.993 & 0.24 & 29.2\\\\ size=50\\%$\\times$7 $''$ & & & & \\\\ NGC 4038 (one) & -2.46 & 0.887 & 0.88 & 17.6 \\\\ NGC 4038 (two) & -6.67 & 0.947 & 1.10 & 6.5\\\\ & -1.45 & 0.995 & 0.49 & 29.9\\\\ size=75\\%$\\times$7 $''$ & & & & \\\\ NGC 4038 (one) & -2.47 & 0.885 & 0.34 & 17.6 \\\\ NGC 4038 (two) & -6.74 & 0.949 & 4.29 & 6.5\\\\ & -1.42 & 0.997 & 1.85 & 30.5\\\\ variable size & & & & \\\\ NGC 4038 (one) & -1.90 & 0.885 & 0.36 & 22.9 \\\\ NGC 4038 (two) & -5.40 & 0.962 & 0.34 & 8.0\\\\ & -1.16 & 0.928 & 0.26 & 37.5\\\\ \\hline NGC 4039 (one) & -3.21 & 0.847 & 0.81 & 13.5 \\\\ NGC 4039 (two) & -11.34 & 1 & 1.01 & 3.8 \\\\ & -1.27  & 1 & 0.55 & 34.1 \\\\ \\hline Overlap (one) & -5.67 & 1 & 0.07 & 7.7\\\\ \\hline M 83 (one)    & -2.73 & 0.999 & 0.45 & 15.9\\\\ \\hline \\end{tabular}  $^{1}$: The total source-averaged CS column density is derived using the following formulae: N(CS)=$10^{intersect} \\times$Q$_{rot}$(T$_{rot}$) where \\textit{intersect} is the value of the intersection of the linear regression with the y-axis (see Figs. \\ref{fig:11} to \\ref{fig:16}) and Q$_{rot}$(T$_{rot}$) is the partition function at the rotational temperature T$_{rot}$ (linear interpolation realized when the available range of T$_{rot}$ was not relevant for the studied case - see text). The rotational temperature is calculated using the formulae: T$_{rot} = - \\frac{1}{slope} \\times log(e)$. \\end{table*}  \\begin{table} \\caption{Results of the LVG model analysis: the physical properties of the best LVG model (having the lowest-$\\chi^{2}$ value) for both the low- and high-temperature gas components. We remind the reader that these values are only indicative since degeneracy in LVG models occurs.}\\label{tab:LVG} \\begin{center} \\begin{tabular}{ccccc} \\hline Source & T$_{K}$ & n(H$_{2}$) & N(CS) \\\\ & in K &  ($\\times 10^{5}$ cm$^{-3}$) & ($\\times 10^{14}$cm$^{-2}$) &\\\\ \\hline NGC 253-1 & 30 & 0.40 & 2.50 \\\\ & 65 & 25.0 & 0.63 \\\\ \\hline NGC 253-2 & 15 & 0.16 & 4.00\\\\ & 70 & 10.0 & 1.00\\\\ \\hline NGC 1068 & 20 & 1.50 & 6.30\\\\ & 65 & 40.0 & 0.63\\\\ \\hline IC 342   & 15 & 0.63 & 1.00 \\\\ & 50 & 16.0 & 0.25 \\\\ \\hline M82-NE   & 10 & 1.00 & 4.00 \\\\ & 65 & 6.30 & 1.00 \\\\ \\hline M82-SW  & 15 & 1.60 & 2.50 \\\\ & 45 & 6.30 & 0.63\\\\ \\hline NGC 4038 & 10 & 1.00 & 0.25 \\\\ & 50 & 40.0 & 0.16 \\\\ \\hline NGC 4039 & 10 & 0.40 & 1.00 \\\\ & 45 & 2.50 & 0.40 \\\\ \\hline Overlap  & 10 & 2.50 & 0.06 \\\\ \\hline M 83     & 65 & 6.30 & 0.15 \\\\ \\hline \\end{tabular} \\end{center} \\end{table}   \\begin{figure} \\includegraphics[height=3cm]{fig1.jpg} \\caption{CS(7-6) spectrum in the center of NGC~253. The angular resolution is shown in the upper right corner. The thick vertical black line symbolizes the V$_{LSR}$ of the source. For NGC~253, we have applied a two-components fit (thin black lines) after having smoothed the entire signal to a velocity resolution of $\\approx 10$ kms$^{-1}$. The fit, on the left is for the emission from NGC~253-1 while on the right, it is for the NGC~253-2 emission.}\\label{fig:1} \\end{figure}   \\begin{figure} \\includegraphics[height=3cm]{fig2.jpg} \\caption{Marginal detection of the CS(7-6) line in the center of NGC~1068 ($\\Delta$v$=13.7$kms$^{-1}$).}\\label{fig:2} \\end{figure}   \\begin{figure} \\includegraphics[height=3cm]{fig3.jpg} \\caption{Marginal detection of the CS(7-6) line in the center of IC~342 ($\\Delta$v$=13.7$kms$^{-1}$).}\\label{fig:3} \\end{figure}  \\begin{figure} \\includegraphics[height=3cm]{fig4.jpg} \\caption{Marginal detection of the CS(7-6) line in the center of Henize~2-10 ($\\Delta$v$=13.7$kms$^{-1}$).}\\label{fig:4} \\end{figure}  \\begin{figure} \\includegraphics[height=14cm]{fig5.jpg} \\caption{CS spectra in the NE molecular lobe of M~82. The velocity resolution from top to bottom is $\\Delta$v$=12.2$kms$^{-1}$, 8.1kms$^{-1}$, 24.5kms$^{-1}$, 9.7kms$^{-1}$ and 13.7kms$^{-1}$, respectively.}\\label{fig:5} \\end{figure}   \\begin{figure} \\includegraphics[height=9cm]{fig6.jpg} \\caption{CS spectra in the SW molecular lobe of M~82. The velocity resolution from top to bottom is $\\Delta$v$=12.2$kms$^{-1}$, 24.5kms$^{-1}$ and 13.7kms$^{-1}$, respectively.}\\label{fig:6} \\end{figure}   \\begin{figure} \\includegraphics[height=18cm]{fig7.jpg} \\caption{CS spectra in NGC~4038 (one of the nuclei in The Antennae Galaxies). The velocity resolution from top to bottom is $\\Delta$v$=24.5$kms$^{-1}$, 16.3kms$^{-1}$, 24.5kms$^{-1}$, 10.0kms$^{-1}$, 10.0kms$^{-1}$ and 13.7 kms$^{-1}$, respectively.}\\label{fig:7} \\end{figure}   \\begin{figure} \\includegraphics[height=12cm]{fig8.jpg} \\caption{CS spectra in NGC~4039 (one of the nuclei in The Antennae Galaxies). The velocity resolution from top to bottom is $\\Delta$v$=24.5$kms$^{-1}$, 16.3kms$^{-1}$, 12.2kms$^{-1}$ and 19.1kms$^{-1}$, respectively.}\\label{fig:8} \\end{figure}   \\begin{figure} \\includegraphics[height=15cm]{fig9.jpg} \\caption{CS spectra in the Overlap position (The Antennae Galaxies). Due to poor weather conditions, no detection is seen neither in the CS(4-3), nor the CS(5-4) nor the CS(7-6) lines. The velocity resolution from top to bottom is $\\Delta$v$=24.5$kms$^{-1}$, 16.3kms$^{-1}$, 24.5kms$^{-1}$, 19.6kms$^{-1}$ and 13.7kms$^{-1}$, respectively.}\\label{fig:9} \\end{figure}  \\begin{figure} \\includegraphics[height=3cm]{fig10.jpg} \\caption{CS(7-6) spectrum in the center of M~83 ($\\Delta$v$=13.7$kms$^{-1}$).}\\label{fig:10} \\end{figure}   \\begin{figure} \\includegraphics[height=10.5cm]{fig11.jpg} \\caption{CS rotational diagrams of NGC~253-1 (top plot) and NGC~253-2 (bottom plot). The linear regression using a single-component is represented by dashed black lines while the two-components regression is represented by solid black lines. Detections are represented by black squares symbols with error bars. These error bars actually correspond to those of the integrated intensities (order of 10-20\\%). When a CS line is marginally detected (upper limit), we represent the corresponding datum with an open white square (see also Figs. \\ref{fig:12}, \\ref{fig:13} and \\ref{fig:15}).}\\label{fig:11} \\end{figure}  \\begin{figure} \\includegraphics[height=5.5cm]{fig12.jpg} \\caption{CS rotational diagram of the center of NGC~1068.}\\label{fig:12} \\end{figure}  \\begin{figure} \\includegraphics[height=5.5cm]{fig13.jpg} \\caption{CS rotational diagram of the center of IC~342.}\\label{fig:13} \\end{figure}  \\begin{figure} \\includegraphics[height=10cm]{fig14.jpg} \\caption{CS rotational diagrams for the M~82-NE (top plot) and the M~82-SW (bottom plot) molecular lobes. }\\label{fig:14} \\end{figure}  \\begin{figure} \\includegraphics[height=14cm]{fig15.jpg} \\caption{CS rotational diagrams of the three positions observed in the Antennae: NGC~4038 (top plot), NGC~4039 (middle plot) and Overlap (bottom plot). For the NGC~4038 CS(5-4) line, we have used a white open triangle and a white open square for representing the IRAM-30m and the JCMT observations, respectively.}\\label{fig:15} \\end{figure}  \\begin{figure} \\includegraphics[height=5cm]{fig16.jpg} \\caption{CS rotational diagram of the center of M~83. }\\label{fig:16} \\end{figure}"
    },
    "0910/0910.3694_arXiv.txt": {
        "abstract": "% ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.1144_arXiv.txt": {
        "abstract": "We study the hadron-quark phase transition with the finite size effects at finite temperature. For the hadron phase, we adopt  a realistic equation of state in the framework of the Brueckner-Hartree-Fock theory including hyperons. The properties of the mixed phase are clarified by considering the finite size effects under the Gibbs conditions. We find that the equation of state becomes softer than that at zero-temperature for some density region. We also find that the equation of state gets closer to that given by the Maxwell construction. Moreover, the number of hyperons is suppressed by the presence of quarks. These are characteristic features of the hadron-quark mixed phase, and should be important for many astrophysical phenomena such as mergers of binary neutron stars. ",
        "introduction": "Introduction}  Nowadays the effects of quark matter on various astrophysical phenomena have been studied extensively. For example, the cooling of compact stars have been studied in Ref.~\\cite{page00, blaschke00, blaschke01, grigorian05}. Other examples include the effects of quark matter on gravitational wave radiation~\\cite{lin06, yas07, abdikamalov08}, neutrino emissions~\\cite{hatsuda87, nakazato08, sagert08}, rotational frequencies~\\cite{burgio03}, and the energy release during the collapse from neutron stars to quark stars~\\cite{yas05, zdunik07}, etc.. In particular, the mechanisms of supernovae and gamma-ray bursts  have not been clearly understood; the QCD phase transition may take place in such phenomena and take critical role~\\cite{gentile93}.  However, there are left many uncertainties for the hadron-quark phase transition, e.g. the equation of state (EOS) of quark matter or deconfinement mechanism. Assuming the quark deconfinement transition to be of first order, we find it causes a thermodynamical instability and the mixed phase appears around the critical density. Since there are two conserved quantities, baryon number and electric charge, the phase equilibrium in the mixed phase must be carefully treated by applying the Gibbs conditions \\cite{glendenning92}, instead of the Maxwell construction. A simple treatment of the mixed phase may be the bulk Gibbs calculation, where phase equilibrium of two bulk matter is considered without electromagnetic interaction and surface tension. Generally the properties of the mixed phase should strongly  depend on electromagnetic interaction and surface tension, and these  effects, sometimes called ``{\\it the finite-size effects}\", lead to the non-uniform \"Pasta\" structures. The EOS of the mixed phase becomes similar to the one under the bulk Gibbs calculation for weak surface tension, and to the one given by the Maxwell construction for strong surface tension~\\cite{voskresensky03,endo06,maruyama07}. The charge screening is also important for their mechanical instability.  In the previous papers these finite-size effects have been properly taken into account to elucidate the properties of the pasta structure and demonstrate the importance of the charge screening at zero temperature \\cite{maruyama07}. However, finite temperature comes in many cases such as relativistic heavy-ion collisions and astrophysical phenomena. In this paper, we study the hadron-quark mixed phase with the finite-size effects at finite temperature by extending the previous works of Maruyama et al.~\\cite{maruyama07}. We adopt the Brueckner-Hartree-Fock EOS by Baldo et al. for the hadron phase~\\cite{baldo98}. The EOS includes hyperons as well as nucleons, interacting with hadronic two-body forces and nucleonic three-body forces. For the quark phase, we adopt the thermodynamic bag model for simplicity. We impose the Gibbs conditions on the phase equilibrium, taking into account the finite-size effects.  This paper is organized as follows. In Sec.~II, we outline our framework. In Sec.~III, we present numerical results. Sec.~IV is devoted to the conclusion and discussion where we give some astrophysical implications of our results. ",
        "conclusions": "We have studied the hadron-quark phase transition at finite temperature. We have taken into account the {\\it finite-size effects}  imposing the Gibbs conditions on the phase equilibrium, and calculating the density profiles in a self-consistent manner.  At finite temperature, EOS of the hadron-quark phase transition gets close to that given by the Maxwell construction. It is due to the mechanical instability of the geometrical structure induced by the thermal effect. Pressure of the mixed phase at finite temperature are 10--30 \\% smaller than that at zero-temperature though the similar behavior appears without hyperons~\\cite{burgio08, yasutake09}. This behavior is characteristic of the hadron-quark mixed phase which can be explained by the large degree of freedom in the quark phase. Hyperon fractions are suppressed by the appearance of the mixed phase, as in the case of zero temperature.  Our calculations are subject to the  neutrino-free~(low $Y_l$) case at finite temperature. Such situation will appear in mergers of neutron star-neutron star binaries or black hole-neutron star binaries~\\cite{shibata06}, and our result may change their dynamical aspects. Of course, we should take into account isentropic and $Y_l\\neq 0$ situation in the core of supernovae. This work is now in progress.  Finally we note again that EOS has many uncertainties, especially for quark matter. We simply adopted the thermodynamic bag model in this paper, and used the density-temperature independent bag constant and surface tension, while it would be interesting to include such dependence for a realistic description~\\cite{baldo03, voskresensky09}. Moreover, chiral restoration or color super conductivities may also change our results~\\cite{kashiwa07,kashiwa07b,jorge07,fukushima08,yasutake09}. These are open questions for astrophysics and nuclear physics."
    },
    "0910/0910.1002_arXiv.txt": {
        "abstract": "I show how to reintroduce velocity dispersion into perturbation theory (PT) calculations of structure in the Universe, i.e., how to go beyond the pressureless fluid approximation, starting from first principles. This addresses a possible deficiency in uses of PT to compute clustering on the weakly non-linear scales that will be critical for probing dark energy. Specifically, I show how to derive a non-negligible value for the (initially tiny) velocity dispersion of dark matter particles, $\\left<\\delta v^2\\right>$, where $\\delta v$ is the deviation of particle velocities from the local bulk flow. The calculation is essentially a renormalization of the homogeneous (zero order) dispersion by fluctuations 1st order in the initial power spectrum. For power law power spectra with $n>-3$, the small-scale fluctuations diverge and significant dispersion can be generated from an arbitrarily small starting value -- the dispersion level is set by an equilibrium between fluctuations generating more dispersion and dispersion suppressing fluctuations. For an $n=-1.4$ power law normalized to match the present non-linear scale, the dispersion would be $\\sim 100~\\kms$. This $n$ corresponds roughly to the slope on the non-linear scale in the real $\\Lambda$CDM Universe, but $\\Lambda$CDM contains much less initial small-scale power -- not enough to bootstrap the small starting dispersion up to a significant value within linear theory (viewed very broadly, structure formation has actually taken place rather suddenly and recently, in spite of the usual ``hierarchical'' description). The next order PT calculation, which I carry out only at an order of magnitude level, should drive the dispersion up into balance with the growing structure, accounting for small dispersion effects seen recently in simulations. ",
        "introduction": "Observing the large-scale density fluctuations in the Universe is one of the best ways we have to approach many fundamental questions about the Universe, e.g., understanding inflation, dark matter, dark energy, the curvature of the Universe, neutrino masses, possible extra dimensions, modifications of gravity, etc. \\citep[e.g.,][]{2009ApJS..180..330K, 2005PhRvD..71j3515S, 2006JCAP...10..014S,2006PhRvL..97s1303S,2007PhRvD..76h3004S, 2003ApJ...594L..71A,2009arXiv0907.2257D,2009arXiv0908.2285L, 2009MNRAS.398..321G,2009arXiv0909.0751B,2008PhRvD..78l3534T, 2008arXiv0810.0323M,2008PhRvD..78l3519M, 2009arXiv0906.4545S, 2009arXiv0906.4548C,2009arXiv0907.5220M}. Statistics of the current density fluctuations can be used to infer statistics of the small initial perturbations from which they grew, and in turn to understand the physics of the very early Universe. Measuring the evolution of large-scale structure (LSS) over time tells us about the present matter content of the Universe and the dynamical rules its evolution follows. Before we can learn anything about fundamental properties of the Universe, however, we must be able to compute the directly observable astrophysical quantities (e.g., the CMB \\cite{2009ApJS..180..296N}, galaxy clustering statistics \\citep{2006PhRvD..74l3507T}, \\lyaf\\ absorption \\citep{2005ApJ...635..761M,2006ApJS..163...80M,2006MNRAS.365..231V, 2007PhRvD..76f3009M,2003ApJ...585...34M,2002ApJ...580...42M}, galaxy ellipticity correlations used to measure weak lensing \\citep{2006ApJ...647..116H}, galaxy cluster/Sunyaev-Zel'dovich effect (SZ) measurements \\citep{2008RPPh...71f6902A}, and possibly future 21cm surveys \\citep{2008PhRvL.100i1303C}) given a hypothetical underlying model.  As observational statistics become more and more precise, with the potential to measure more and more subtle differences between models, the requirements on the phenomenological theory needed for their interpretation become correspondingly more stringent.  Currently, linear-order perturbation theory \\citep{1980lssu.book.....P, 1996ApJ...469..437S,1984ApJ...284L...9K,1987MNRAS.227....1K} provides our primary means of calculating LSS observables for cosmological parameter estimation and model testing, with only ad hoc, and recently demonstrably inadequate, attempts to correct for non-linearity once it becomes non-negligible \\citep{2008MNRAS.385..830S,2006PhRvD..74l3507T,2006JCAP...10..014S, 2008JCAP...07..017H,2009ApJS..180..330K} (a vast number of papers have been written on beyond-linear calculations, but most of these are never used in parameter estimation papers).  Linear theory can robustly describe observations on very large scales or at very early times, but breaks down when the perturbations become too large on a given scale. When considering gravitational evolution only (i.e., dark matter only), numerical simulations can be used to compute the fully non-linear evolution of the density field to high accuracy (with a lot of care and computer power \\cite{2008CS&D....1a5003H,2008arXiv0812.1052H,2009arXiv0902.0429H}), however, as discussed extensively in \\cite{2009JCAP...08..020M}, numerical simulations are a less than ideal tool for interpreting future precision observations, once one considers real observables which are inevitably influenced by baryonic effects (e.g., star formation, and the accompanying complication of general gas dynamics). To summarize the argument in \\cite{2009JCAP...08..020M}: Beyond linear order perturbation theory (PT) \\citep{1980lssu.book.....P,1981MNRAS.197..931J,1983MNRAS.203..345V, 1984ApJ...279..499F,1986ApJ...311....6G,1996ApJ...473..620S, 2002PhR...367....1B,2006ApJ...651..619J} should provide the primary means of interpreting very high precision future LSS data, just like linear theory has provided the primary means in the past. The range of scales over which higher order PT will be necessary (i.e., linear theory is insufficient) and sufficient (i.e., it will be accurate enough after computing a modest number of terms) will become larger as the precision of observations increases, while the chances of robustly, completely, convincingly describing the full precision of the observations using inevitably somewhat ad hoc prescriptions for star formation and related things in simulations becomes more remote. The key difference between PT and simulations is the fact that perturbative clustering can be completely described by a finite set of well-defined parameters, no matter how complicated the small-scale physics is (at least as conjectured in \\cite{2009JCAP...08..020M}), while the need for fully non-linear calculations implies that there is generally no bound on the number or form of free parameters (more or less by definition).  The idea that the importance of PT relative to simulations is increasing with time may seem backwards relative to conventional wisdom, however, my argument is that this conventional wisdom was developed for the era of not very precise observations, when corrections to linear theory were already too large to be described by PT by the time they were large enough to be statistically significant, i.e., past PT work was premature.  To be clear, I am not saying that simulations will not be extremely useful, only that they are unlikely to be the leading tool for extracting fundamental cosmological information from very high precision observations (much like the situation in high energy physics, where lattice QCD simulations provide much qualitative, and recently even quite precise quantitative, insight \\cite{2008Sci...322.1224D}, but high precision constraints on models are made primarily in the regime accessible to perturbation theory \\cite{2006JPhG...33....1Y}).  Following the above line of reasoning, this paper is aimed at building up our ability to do calculations based on perturbation theory. It deals exclusively with the gravity-only part of the calculation, however, this should be viewed only as an intermediate goal. PT is most necessary for practical computations of observables which can not realistically be done from first principles in a simulation -- understanding how to do computations for dark matter-only is simply a prerequisite for computing these observables. This continues a recent line of work related to renormalization/resummation methods \\cite{2006PhRvD..73f3520C, 2006PhRvD..73f3519C,2007JCAP...06..026M,2007PhRvD..76h3517I, 2008ApJ...674..617T,2007PhRvD..75d3514M,2008A&A...484...79V, 2008PhRvD..77f3530M,2008MPLA...23...25M,2008PhRvD..77b3533C, 2008JCAP...10..036P, 2008PhRvD..78j3521B,2008PhRvD..78h3503B, 2009JCAP...06..017L,2009MNRAS.397.1275W,2009PhRvD..80d3531C, 2009PASJ...61..321N}. In my opinion, the key to maximizing the usefulness of all of this work is eventually connecting it to galaxy biasing and related complications \\cite{2005PhRvD..71d3511S,2006PhRvD..74j3512M, 2008PhRvD..78h3519M,2009JCAP...08..020M}.  One traditional potential limitation of LSS perturbation theory, which I focus on addressing in this paper, is that it assumes the particles at a point in space all have exactly the same velocity \\cite{2009PhRvD..80d3504P}. The equations used are those of a pressureless fluid. This is often referred to as the ``single-stream approximation'', however, I will avoid this language as it assumes a certain picture of large-scale structure formation that may or may not have anything to do with reality. I say this because, for typically observed times, stream crossing is actually ubiquitous -- in fact, there are many orders of magnitude in scale of deeply non-linear structure, to the point where the initial conditions on very small scales may be effectively forgotten. Of course, the standard language implicitly means no stream crossing when the field is in some sense smoothed on the typical scale where the perturbation theory is supposed to apply. When you look at it this way, however, it is clear that, if coldness is a good effective theory, it is not simply because the particles were initially cold, there is an additional requirement that the velocity field remains effectively cold on scales just smaller than the ones of interest.  To give a qualitative preview of the paper: The exact equation for the evolution of collisionless particles is the Vlasov equation for the full distribution function in 6-dimensional phase space (particle density in position and momentum space). The standard hydrodynamic equations are derived by taking moments of the Vlasov equation with respect to momentum -- the zeroth moment gives an equation for the evolution of density, the first moment gives an equation for the evolution of bulk velocity, and higher moments are normally dropped, e.g., the second moment which would describe velocity dispersion. At first glance, one might think that standard PT could be extended by simply adding the evolution equation for the 2nd and possibly higher moments, however, that turns out to be less straightforward than it sounds. Viewed conventionally, dropping these higher moments is not really an added approximation in standard Eulerian perturbation theory. The lowest order evolution equations contain only a Hubble drag term, so any small initial velocity dispersion will rapidly become even smaller. Higher order equations contain no source terms, meaning that the higher order results are always proportional to the tiny starting value. \\cite{2001A&A...379....8V} showed that even perturbation theory using the full distribution function directly leads to the same result. I will show, however, that this argument for dropping dispersion from the perturbation theory is faulty, in that, while the lowest order terms are very small, the series is very rapidly diverging, in the sense that higher order terms are increasing in size instead of decreasing. This implies that some rearrangement of the series is needed, as in a resummation or renormalization group calculation. In case this mathematical reasoning is not sufficient motivation, recently \\cite{2009PhRvD..80d3504P} computed the velocity dispersion generated in N-body simulations, finding it to be small but not completely dynamically negligible.  There has been much discussion in the past about different ways of adding velocity dispersion, or more general changes to the small-scale effective theory used for PT calculations \\cite[e.g.,][]{2004PhRvD..70f4010T,2005PhRvD..71f7302R, 2006PhRvD..73b4024S,2002MNRAS.330..907M,2005A&A...438..443B, 2007JPhA...40.6849G,2009ApJ...700..705S}, however, these papers all lacked a first principles derivation, starting from the exact equations, of the model they use. This made their usefulness for precision calculations questionable, as there were always added assumptions and/or free parameters. The point of this paper is to show how to do a straightforward computation that takes us directly from the initial homogeneous, cold, starting theory to a theory with highly developed, potentially hot, small-scale structure.  The rest of the paper is as follows: In \\S\\ref{seccalc} I show how to use renormalization group-inspired ideas to reintroduce velocity dispersion from first principles. In \\S\\ref{secredshiftspace} I give a short discussion of the implications of these velocity dispersion calculations for redshift-space distortions. Finally, in \\S\\ref{secdiscuss} I discuss the results. ",
        "conclusions": "}  The deficiencies of standard perturbation theory are a lack of control of the higher moments of the velocity distribution function, beyond the pressureless fluid approximation, problems with accuracy related to the fact that PT integrals include small scales that are generally highly non-linear, even when the scales we are interested in are weakly non-linear, and, of course, simply not working on small scales where the fluctuations are large. The deficiencies of numerical simulations are speed, and the related fact that statistics (e.g., the power spectrum) cannot be computed directly, but only as averages over realizations of the random density field, which must contain many orders of magnitude more degrees of freedom (e.g., a billion) than one really needs to describe the statistic of interest (e.g., a few dozen for the power spectrum). Additionally, and probably most importantly, the cumbersome, opaque nature of simulation results is greatly exacerbated when they are used to model galaxies or other observables instead of just dark matter, while the advantage of being more or less exact for gravity alone (at least in the straightforward limits of large box size and particle density) is lost.  This paper directly addresses the PT deficiency of missing higher moments of the velocity distribution function, showing how they can be re-activated and generated at a significant level, starting from first principles. The approach here may also improve PT by providing natural regulation of small-scale structure, i.e., the Jeans smoothing that arises here has the same kind of effect as the propagator resummation of \\cite{2006PhRvD..73f3520C,2006PhRvD..73f3519C}. The philosophy of this paper is that the cold ``streams'' often discussed as ``crossing'' are mythical objects -- what one sees in the real Universe is always some evolution of effectively warm (although maybe not very warm) material. The meaning of this idea is very clear for power law power spectra, where, as we saw, non-trivial effects of velocity dispersion can be computed without any initial dispersion entering the discussion. The dispersion bootstraps itself up from an arbitrarily small start. The temperature is locked into a sort of equilibrium with the growth of structure. The situation is not as clear for $\\Lambda$CDM power spectra, not so much because the $\\Lambda$CDM is cold as because there is very little small-scale power in these models, so the dispersion computed to linear order in the power spectrum remains well below the non-linear scale (although orders of magnitude larger than it would be if there was no structure). This situation will change when calculations are done to higher order, where I showed that there is enough small-scale power generated to produce dispersion that would be far too large in absence of feedback on the structure formation itself. The initial velocity dispersion should be rendered irrelevant when the system moves into a sort of self-regulating mode, like the power law example, where the velocity dispersion and growth of structure are tightly coupled by feedback between them. Ultimately, it seems likely that the best effective small-scale model for doing perturbation theory will involve some more general balance between different pieces, i.e., density, velocity divergence, dispersion, and maybe even vorticity, etc., because we know that physically this is what the small-scale structure really is, i.e., halos which can only be described as a delicate balance of the original perturbative LSS quantities. To put this another way: hopefully, when all of the relevant elements are included, there will  be some fixed-point structure for small scales with a clear physical interpretation and effectively far fewer than the original number of degrees of freedom (akin to the fixed-point power law found in \\cite{2007PhRvD..75d3514M}, but with richer structure).  This paper does not exactly contain useful quantitative take-home results. The results are primarily a procedure to follow for future calculations. If there is a single equation that best represents the results, it is probably Eq. (\\ref{eqkNLsigRG}), which shows how, for a power law power spectrum, the velocity dispersion tracks the non-linear scale, with Jeans filtering erasing more and more small-scale structure as the larger scale structure grows. Eqs. (\\ref{eqTnonlinear}) and (\\ref{eqtsigevolution}) are also critical, in that they show generally how homogeneous velocity dispersion is generated out of fluctuations. The key new variable, equivalent to $\\delta$ and $\\theta$, but for perturbations in velocity dispersion, is $\\pi \\propto \\partial_i \\partial_j \\sigma^{ij}$.  The next step in this line of work is to derive to the next-order equations like Eq. (\\ref{eqPddevolution}-\\ref{eqPppevolution}) (but including bispectrum equations), which are needed to have any chance of properly describing $\\Lambda$CDM. Then the results can be tested by comparison to numerical simulations. While $\\Lambda$CDM is of course the ultimate goal, tests of the theoretical concepts here might be more conveniently done with power law simulations, particularly ones with relatively blue spectra, as in \\cite{2009MNRAS.397.1275W}. Any of the renormalization approaches that have been developed recently can probably be applied to the problem of renormalizing the velocity dispersion, at least in principle, i.e., resumming a set of terms that generates the dispersion, either explicitly or through a renormalization group equation. The obstacle to doing this elegantly may be the difficulty of obtaining analytic solutions as the evolution equations become more complicated (this problem of requiring analytic solutions is, I think, the primary reason to favor the ``numerical evaluation of evolution equations for statistics'' approach advocated in this paper over other, more completely analytic, recent approaches).  One might ask ``why stop with the 2nd moment of the distribution function, i.e., why can we drop $q_{ijk}$ from Eq. (\\ref{eqdispersion})?'' One hope, which will need to be verified by future calculations, is that the effect of increasingly high moments on large scales may take the form of a gradient expansion, i.e., in Fourier space a series where the effects of increasingly high moments enter multiplied by increasing powers of $k$.  In this case, we would expect that, on scales where the effect of the 2nd moment are already small, the effects of higher moments will be even smaller.  The result for redshift-space space distortions (Eq. \\ref{eqredshiftspace}) leads to the question:  Why do we use $\\vv$, and in this paper $\\sigma_{i j}$, as the variables to be solved for perturbatively, rather than, e.g., momentum $\\left(1+\\delta\\right)\\vv$ and kinetic energy $\\left(1+\\delta\\right)\\left(v_i v_j+\\sigma_{ij}\\right)$? An answer one might have considered was that velocities are needed to compute redshift-space distortions, but here we see that the most direct quantities for that are in fact momentum and energy, with additional perturbative calculations needed when starting from $\\delta$, $\\vv$, and $\\sigma_{ij}$. A change to total kinetic energy would have the potentially substantial effect of increasing the Jeans-like smoothing of the linear power, because the zero order pressure would be larger.  I do not see any clear a priori reason to favor one option over another -- they are just different ways of arranging a series of terms, which should be equivalent if one could include an infinite number of terms.  Note that, while the choice of variables is optional, the renormalization of the energy-related variable is not optional -- it simply makes no sense to ignore the fact that the size of terms in the series describing one of your variables is increasing rapidly, rather than decreasing, with order. There is a lot of circumstantial evidence that using total kinetic energy could be useful.  The renormalization/resummation of the propagator in \\citep{2006PhRvD..73f3520C} leads to filtering much like the Jeans filtering we find due to velocity dispersion, but with scale given by the bulk velocity power spectrum.  The Lagrangian PT-based scheme of \\citep{2008PhRvD..77f3530M} leads to a similar result. Another possibility to consider would be the evolution of large-scale fields with the small-scale structure explicitly averaged out, which would naturally lead to the inclusion of small-scale bulk velocities in the dispersion term \\cite{2000PhRvD..62j3501D,2002MNRAS.334..435D,2003AN....324..560D}. One might also consider using the Schr\\\"{o}dinger equation representation of the exact Vlasov equation, proposed in \\cite{1993ApJ...416L..71W}, combined with the approach of this paper.  In the end, one could view the approach in this paper as a re-activation of the program under discussion in papers like \\cite{1977ApJS...34..425D}, which set out to integrate the BBGKY equations numerically.  This reconsideration is timely because of several developments over the last thirty years.  We now know the appropriate class of models to focus on, especially including the appropriate the initial conditions. There has been a lot of work on both perturbation theory beyond linear order, and on N-body simulations, with the limitations of each teaching us a lot about what is needed from new methods. We also have a specific calculation to focus on:  the scales where baryonic acoustic oscillations (BAO) are observable \\citep{1998ApJ...496..605E, 2003ApJ...594..665B,2003PhRvD..68h3504L,2003ApJ...598..720S, 2004ApJ...615..573M,2005ApJ...631....1G, 2005MNRAS.357..429A,2005MNRAS.363.1329B,2006MNRAS.365..255B, 2007PhRvD..76f3009M}, which points us toward the perturbative approach that motivates truncating the hierarchy and believing that high precision can be achieved (in contrast to \\cite{1977ApJS...34..425D}, who were focused on the more strongly non-linear regime, where the truncation used here is not well-motivated). The same scales are also potentially the most powerful for other measurements based on, e.g., redshift-space distortions \\cite{2008arXiv0810.0323M}.  The basic form of calculation I do here is completely standard in some other areas of cosmology.  For example, the evolution of the homogeneous (background) value of an interacting scalar field in the early Universe is affected by quantum and thermal fluctuations.  Its evolution is not described by the original equation of motion with all perturbations ignored, but by a renormalized effective potential, which is at least formally infinitely different from the original potential, i.e., completely dominated by the part due to fluctuations \\cite{2005pfc..book.....M}. Another interestingly similar calculation is the development of equilibrium after preheating after inflation. \\cite{2000hep.ph...11159F, 2001PhRvD..63j3503F, 2006JCAP...07..006D,2006PhRvD..73b3501P,2007PhRvD..75d3518F} perform fully non-linear simulations of the interaction of scalar fields, much like large-scale structure simulations, with the added twist that the background density and pressure are affected by the fluctuations.  The evolution of the scale factor is calculated by taking spatial averages over the fluctuations as the simulation is running.  In the beginning, the nearly homogeneous inflaton field dominates, and a very naive calculation might compute the expansion of the simulation in advance assuming homogeneous evolution, but by the end of the simulation the homogeneous component has actually practically disappeared, with the background evolution completely dominated by the fluctuations, which behave like radiation. Of course no one would ever do the very naive calculation just mentioned, where the affect of the fluctuations on the background equation of state is ignored... except that this is what we do when we assume that the tiny initial temperature of CDM means that it will forever remain cold.  Substantial work remains to determine if the approach in this paper enhances the accuracy of predictions of quasi-linear clustering in the real Universe; however, it is now possible to consider an effect that was previously outside of any first-principles computational control in this kind of perturbation theory. The bottom line of this paper is that even linear theory fluctuations necessarily imply a significant one-loop renormalization of the background velocity dispersion."
    },
    "0910/0910.3836_arXiv.txt": {
        "abstract": "We explore potential strategies for testing General Relativity via the coherent motions of galaxies. Our position at $z=0$ provides the reference point for distance measures in cosmology. By contrast, the Cosmic Microwave Background at $z \\simeq 1100$ acts as the point of reference for the growth of large scale structure. As a result, we find there is a lack of synergy between growth and distance measures. We show that when measuring the gravitational growth index $\\gamma$ using redshift-space distortions, typically $80\\%$ of the signal corresponds to the local growth rate at the galaxy bin location, while the remaining fraction is determined by its behaviour at higher redshifts.  In order to clarify whether modified gravity may be responsible for the dark energy phenomenon, the aim is to search for a modification to the growth of structure. One might expect the magnitude of this deviation to be commensurate with the apparent dark energy density $\\Omega_\\Lambda(z)$. This provides an incentive to study redshift-space distortions (RSD) at as \\emph{low} a redshift as is practical. Specifically, we find the region around $z = 0.5$ offers the optimal balance of available volume and signal strength. ",
        "introduction": "Dark energy is a low-redshift phenomenon, and in accordance with the standard $\\Lambda CDM$ model, it appears to exert a rapidly decaying influence towards higher redshifts, at a rate approaching $(1+z)^{-4}$. Physical models of dark energy may be distinguished by an equation of state $w \\neq -1$, while a break from general relativity would likely exhibit a distinctive structure formation history. The consequences of a modification to general relativity are somewhat speculative at this stage, but naturally one might expect any alteration of the growth rate to become particularly prominent at late times, in accordance with the observed change in global dynamics. For instance, the $f(R)$ models explored by Hu \\& Sawicki \\cite{2007PhRvD..76j4043H} found this to be the case on large scales, as did He et al \\cite{2009JCAP...07..030H} with their interacting model.  A number of probes are capable of measuring $w$ via its influence on the redshift-distance relation. It is rather more difficult to study the growth rate, due to our uncertainty in the behaviour of galaxy bias, but there are currently two promising avenues available for future exploration. Weak gravitational lensing provides a direct measurement of the dark matter distribution, and its evolution with redshift. However the focus of this work will be redshift-space distortions, which exploit the relationship between the large-scale coherent velocities of galaxies and the growth rate of perturbations.  Weight functions have previously been applied to the equation of state as a function of redshift $w(z)$ \\cite{2003MNRAS.343..533D,2005PhRvD..71h3501S,2006PhRvD..73h3001S}, and this work extends the concept to the growth index $\\gamma$ \\cite{2005PhRvD..72d3529L}. They are designed to illustrate the redshift sensitivity of a given survey, and this is often quite different from the source redshift distribution. Weight functions are closely related to, and may be derived from, the principal component approach. In related work, Zhao et al. \\cite{2009arXiv0905.1326Z} recently explored the principal components of the metric ratio, a different quantity also often denoted by $\\gamma$.  In \\S\\ref{sec:wts} we briefly review the concept of principal components and their application in deriving the weight function. \\S\\ref{sec:rsd} explores the RSD redshift sensitivity to the growth index $\\gamma$, and compares these weight functions to those of the dark energy equation of state. \\S\\ref{sec:opt} addresses the question of which redshift is most efficient at measuring $\\gamma$, while in \\S\\ref{sec:gro} we consider a change of parameterisation to the growth rate $f$. ",
        "conclusions": "At present there are two key approaches to study the evolution of dark matter perturbations, namely redshift-space distortions and weak gravitational lensing. They will provide a crucial piece of evidence in determining whether the phenomenon of dark energy may be attributed to a physical entity, or is simply due to a misunderstanding of the laws of gravity.  In this work we have quantified the epoch at which a galaxy redshift survey would be sensitive to the growth index $\\gamma$. Specifically, given the absence of any weight at redshifts lower than that of the survey, low redshift surveys are left with the significant advantage of probing a broader behaviour $\\gamma(z)$. Another interesting feature we have highlighted is that approximately $80 \\%$ of the ``weight\" for $\\gamma$ for any given redshift bin corresponds to the local value of $\\gamma(z)$ at the location of the bin.  The main limitation of a low redshift survey would be the available comoving volume, along with the stronger prevalence of nonlinear perturbations which may hinder an accurate determination of $\\beta$. However, for a given error on $\\beta$, one reaches a much better determination of $\\gamma$ compared to higher redshifts.  By contrast, a weak lensing survey with source galaxies at any redshift will still be sensitive to the value of $\\gamma(z)$ across \\emph{all} redshifts $(0<z<1100)$. Although this may be restricted at very low redshift by the lensing efficiency.  For a purely cosmic-variance limited redshift survey, it appears the shell surrounding $z \\sim 0.5$ is optimal in its efficiency. Incorporating more realistic constraints, such as limited observing time, will likely serve to lower this value further unless non-linearities prove highly problematic. This preference towards a low-redshift measurement reflects the increased likelihood of discovering a modified gravity signature when searching within the era of dark energy.   \\noindent{\\bf Acknowledgements} \\\\ The author would like to thank S. Bridle, A. Heavens, J. Peacock and P. Norberg  for helpful comments. FS is supported by an STFC rolling grant, and acknowledges the hospitality of the Aspen Center for Physics where part of this work was undertaken."
    },
    "0910/0910.1833_arXiv.txt": {
        "abstract": "The galaxy power spectrum contains information on the growth of structure, the growth rate through redshift space distortions, and the cosmic expansion through baryon acoustic oscillation features. We study the ability of two proposed experiments, BigBOSS and JDEM-PS, to test the cosmological model and general relativity. We quantify the latter result in terms of the gravitational growth index $\\gamma$, whose value in general relativity is $\\gamma\\approx 0.55$. Significant deviations from this value could indicate new physics beyond the standard model of cosmology.  The results show that BigBOSS (JDEM-PS) would be capable of measuring $\\gamma$ with an uncertainty $\\sigma(\\gamma) = 0.043$ (0.054), which tightens to $\\sigma(\\gamma) = 0.031$ (0.038) if we include Stage III data priors, marginalizing over neutrino mass, time varying dark energy equation of state, and other parameters. For all dark energy parameters and related figures of merit the two experiments give comparable results.  We also carry out some studies of the influence of redshift range, resolution, treatment of nonlinearities, and bias evolution to enable further improvement. ",
        "introduction": "Surveys of large-scale structure in the universe provide a rich resource for testing our understanding of cosmology. Future surveys will cover nearly the full sky to redshifts far deeper than are currently studied, mapping out some 10 billion years of history.  The great statistical power and leverage from depth will allow detailed examination of the cosmological framework by carrying out a simultaneous fit of a substantial suite of relevant parameters. One particularly attractive prospect is the capability to put to the test the predictions of Einstein gravity for the growth of structure and its consistency with the cosmic expansion history.  We consider next-generation surveys mapping the distribution of galaxies in three dimensions to redshifts of order $z=2$. A goal of this study is to determine the capabilities of such surveys. In particular we aim to estimate realistic constraints from a global parameter fit on the gravitational growth index $\\gamma$, which can characterize deviations from general relativity. The second goal is to examine how the survey characteristics such as redshift range, resolution, and galaxy selection affect those capabilities.  In Sec.~\\ref{smethod} we review the formalism for extracting cosmological information from galaxy correlation measurements in terms of the matter power spectrum, and discuss the anisotropic distortion due to measuring in redshift space (rather than position space).  We discuss the relevant set of cosmological parameters in Sec.~\\ref{ssolve} and their influence on the matter power spectrum. The results are analyzed with emphasis on the role of degeneracies between factors that influence growth, including the gravitational growth index, the dark energy equation of state, and neutrino mass. In Sec.~\\ref{ssurvey} we turn to astrophysical and survey characteristics and analyze the effect of the bias level of the selected galaxy populations, the form of the small-scale velocity damping, the spectroscopic survey redshift resolution, and the redshift range of the survey.  This allows quantitative comparison of the capabilities of next-generation (Stage IV) experiments from both ground and space, as well as nearer term (Stage III) experiments.  We conclude in Sec.~\\ref{sconcl} with a summary of the prospects for testing the standard cosmology and revealing clues to dark energy or the breakdown of Einstein gravity. ",
        "conclusions": "The three-dimensional distribution of large scale structure contains information on both the cosmological parameters and testing gravity. We have studied the capabilities of next-generation power spectrum experiments from the ground, BigBOSS, and from space, JDEM-PS, to use the baryon acoustic oscillations, power spectrum shape, and redshift space distortions to test standard cosmology.  The main conclusion is that the two experiments could achieve comparable constraints.  We emphasized the importance of including simultaneously the parameters that affect growth -- the gravitational growth index characterizing deviations from general relativity, the dark energy equation of state value and its time variation, and neutrino mass.  Including these and other cosmological parameters we estimate the uncertainty on determination of the gravitational growth index to be 0.043 for BigBOSS, 0.054 for JDEM-PS, or 0.031 and 0.038 respectively when combined with nearer-term, Stage III experiments.  This represents nearly an order of magnitude improvement over Stage III knowledge.  We have also studied the survey characteristics and confirm that the power spectrum at redshifts $z\\lesssim1$ has strong leverage. This makes the luminous red galaxy component of the survey quite important.  Furthermore, our results demonstrate that shifting the redshift range of the emission line galaxy survey of BigBOSS from $z=1-2$ to $z=0.7-1.7$ can improve the constraints, while adding benefits such as reduced technical complexity and line confusion and increased signal-to-noise and the ability to crosscorrelate galaxy populations of different biases.  Lyman-$\\alpha$ forest spectra from BigBOSS quasars at $z>2$, which we have neglected, will further advance the determination of cosmological parameters.  The prospects for testing standard cosmology and in particular general relativity are promising.  Improved understanding of the translinear density regime and velocities would further extend the number of usable power spectrum modes, while complementarity with other Stage IV experiments utilizing supernova distances, CMB measurements, and weak lensing data would give powerful leverage on both the gravitational growth index and other cosmological parameters.  The capability of probing beyond-Einstein gravity opens up a new window for our understanding of cosmic acceleration and fundamental physics."
    },
    "0910/0910.2417_arXiv.txt": {
        "abstract": "The problem of detecting dark matter filaments in the cosmic web is considered. Weak lensing is an ideal probe of dark matter, and therefore forms the basis of particularly promising detection methods. We consider and develop a number of weak lensing techniques that could be used to detect filaments in individual or stacked cluster fields, and apply them to synthetic lensing data sets in the fields of clusters from the Millennium Simulation. These techniques are multipole moments of the shear and convergence, mass reconstruction, and parameterized fits to filament mass profiles using a Markov Chain Monte Carlo approach. In particular, two new filament detection techniques are explored (multipole shear filters and Markov Chain Monte Carlo mass profile fits), and we outline the quality of data required to be able to identify and quantify filament profiles. We also consider the effects of large scale structure on filament detection. We conclude that using these techniques, there will be realistic prospects of detecting filaments in data from future space-based missions.  The methods presented in this paper will be of great use in the identification of dark matter filaments in future surveys. ",
        "introduction": "In a universe dominated by Cold Dark Matter (CDM) such as our concordance cosmology, $\\Lambda$CDM, hierarchical structure formation takes place with the growth of collapsed objects progressing via merging of smaller objects. Simulations of structure formation show that galaxy clusters are located at the intersections of filaments, which form a web-like structure throughout the universe. The processes of merging and continuous accretion of surrounding matter are thought to occur highly anisotropically,  with the filaments channeling matter along preferred directions. Modern galaxy surveys such as 2dFGRS (e.g. \\cite{Coll01}) and SDSS (e.g. \\cite{yorksdss}, \\cite{Adel08}) have provided a dramatic picture of this so-called `cosmic web', where voids lie between filamentary arms traced by galaxies. From a theoretical perspective, \\citet{Bond96} developed a theory for the cosmic web in a CDM cosmogony, where filaments are a product of a primordial tidal field evolving under non-linear effects.  Detection of inter-cluster filaments has important implications for cosmology, confirming the prediction that structure grows highly anisotropically from small initial perturbations in the cosmic density field. Their detection provides an important validation of our picture of hierarchical structure formation and cluster evolution \\citep{Colfil}.  Dark matter filaments also have a significant role to play when considering the total mass budget of the universe - the galaxies and gas channeled along these filaments may account for as much as half of the baryonic mass budget in the universe.  Inter-cluster filaments have been the target of a number of observational searches. The earliest of these used bremsstrahlung X-ray emission to detect hot gas being channeled along filaments. For example, \\cite{Briel95} used X-ray data from the ROSAT All-Sky Survey, but failed to find evidence for filaments. Subsequent X-ray searches have also proved inconclusive (\\cite{scharf00}, \\cite{KB99}, \\cite{durret03}), the main issue being that it is difficult to ascertain whether strong X-ray emission is due to matter in a filament or due to the clusters interacting. Another proposed observational method looks for overdensities of galaxies relative to the background. \\cite{PD04} reported finding a short filament between galaxy clusters A1079/1084 using such a method. \\cite{Ebeling} report the detection of a filament leading into MACS J0717.5+3745. However, this approach is susceptible to large errors and requires spectroscopic data to establish the redshifts of the target galaxies. An alternative detection technique uses a `skeleton', which is defined as the subset of critical lines joining the saddle points of a field to its maxima while following the gradient's direction. The skeleton formalism has been used in numerical simulations to extract and analyze the filamentary structure of the universe (\\cite{Nov}, \\cite{sousbie}).  Work by \\cite{Hahn} and \\cite{Aragon} demonstrates that it is possible to use the Hessian of either the potential or density field to provide a dynamical classification of filamentary structure. This work was extended by \\cite{romero}. The number of eigenvalues with magnitude above a certain threshold at each grid point determines whether that point belongs to a void, sheet or filament. \\cite{Hahn} investigated the effect filaments have on the properties of dark matter halos, and found that filaments influence the magnitude and direction of halo spin as well as halo shape. There are a number of other filament detection methods, not elaborated here, which are discussed in \\cite{Pim05}.  One of the most promising methods of filament detection is weak gravitational lensing. Since gravitational lensing is insensitive to whether the matter is luminous or dark, and to its dynamical state, it is an ideal probe of dark matter. A number of attempts have been made to image filaments using weak lensing. \\citet{Gray02} claimed to detect a filament in the A901/902 supercluster, but this was later shown by \\cite{Heymans08} to be an artifact of the mass reconstruction. \\cite{Kaiser98} reported a detection of a filament between two of the galaxy clusters in the supercluster MS 0302+17, but these results could not be reproduced by \\cite{gav04}, and hence this detection is called into question. \\cite{Dietrich} used a variety of methods to detect a filament between A222/223 - however, they point out that it is difficult to objectively define whether this lensing signal is caused by a real filament. The Hubble Space Telescope Cosmic Evolution survey (COSMOS) has provided corroboratory evidence for a cosmic network of filaments \\citep{Massey}.  Although there have been a number of observational attempts to detect the filamentary cosmic web with various degrees of success, conclusive evidence has proved elusive in all but a few cases, and widely applicable methods are still lacking. In this paper we assess existing techniques to detect filaments in weak lensing data and present two new, more effective, detection methods. In particular we present a method based on multipole filtering of the shear field, and a method based on Markov Chain Monte Carlo techniques that allows both detection and quantification of filament profiles. Motivated by the promise of weak lensing data from many thousands of galaxy clusters from large surveys such as the Dark Energy Survey (DES) and Large Synoptic Survey Telescope (LSST) we also consider detection in stacked data sets from a number of clusters. These methods are then applied to single and stacked synthetic weak lensing data sets from the fields of clusters from the Millennium Simulation \\citep{springel05}. We will show that weak lensing is an ideal and versatile tool to detect dark matter filaments.  This paper is organized as follows. After a review of the basic lensing formalism (section 2), in section 3 we discuss the various methods by which we may detect filaments using weak lensing. These methods are then applied to data sets from the fields of simulated clusters. In section 4 we describe these simulations and in section 5 we detail our results. In section 6 we analyze the effects of large scale structure noise on our results. Finally, in section 7 we present our discussions and conclusions. ",
        "conclusions": "The typical strength of an intercluster filament is $\\kappa \\sim 0.01$, compared to a convergence in the central cluster regions of $\\kappa \\sim 0.5$. Therefore, the reason filaments are so difficult to detect using any method is due to the fact that they are much weaker than the cluster itself, and are quickly `lost' in the noise originating from the intrinsic ellipticity dispersion of the background galaxies. In this paper, we have shown several methods that can be used to reliably detect filaments. By implementing a variety of methods, we have also shown how versatile weak gravitational lensing is as a tool for detecting filaments - each method described in the paper has different strengths and weaknesses that make them suited to different data sets.  The mass reconstruction is one of the most common methods to detect overdensities in a mass distribution. We applied the finite field mass reconstruction algorithm of \\cite{seitzrec} to both single and stacked clusters. For single clusters, reconstructions using a background galaxy density of 30 arcmin$^{-2}$ weakly detect only the strongest filaments, with the majority of filaments undetected. As background galaxy density increases, so does the ease of filament detection - at 100 arcmin$^{-2}$ we are able to identify some filaments in reconstructions. The mass reconstruction is most effective when used with stacked data sets. The stacking method has an impact on the quality of the results, with the double-cluster method proving more effective than aligning cluster major axes. A major disadvantage of mass reconstructions is that they can give false filamentary structures due to the smearing effect of the smoothing scale. This is not a problem unique to mass reconstructions, and any method which smoothes results over certain scales in an aperture will suffer from the same difficulty. This is a particular problem for closely-spaced clusters, which will appear to merge in the reconstructed image if the smoothing scale is greater than the inter-cluster separation. The second major disadvantage of a mass reconstruction is that it does not allow easy calculation of the significance of any filament detection. The origin of this effect lies in the fact that in any reconstruction, the shear field must be smoothed to avoid infinite noise \\citep{ks93}. This leads to correlated errors in the resulting convergence map, making error bars very difficult to construct. Thirdly, because filaments are weak, they are often located in regions that are dominated by noise. As \\cite{Dietrich} points out, in such a scenario the value of the mass sheet degeneracy can fluctuate by as much as the filament strength between different galaxy populations, making any filament detection in an individual cluster more questionable. This third point does not apply for the stacked clusters, as the stacking ensures that noise is reduced to sub-dominant levels even in filamentary regions.  Therefore, it is advisable to use the mass reconstruction in conjunction with other methods described in this paper, both to confirm any detection and to allow the significance of any detection to be accurately assessed.  In general, the multipole moments of the convergence peformed poorly, and were unable to register any strong filament detections in single clusters. Although \\cite{Dietrich} used $Q^{(2)}$ to provide evidence for a filament in A222/223, this cluster pair was very closely spaced and the corresponding filament was thus much stronger than any considered in this paper. It is also unclear how much of a contribution the filament makes to the quadrupole signal, as it is likely that the majority of the signal comes from the clusters themselves. The quadrupole signal fails to detect filaments in the stacked data due to misalignments. Misalignment between filaments will create a `mass-sheet' in the resulting stacked data, which will give zero quadrupole signal.  It is not only possible to define multipole moments of the convergence, but also of the shear. Here we have shown that the shear monopole is very effective in picking out filaments. The reason for this is that the shear monopole performs a (weighted) sum of all shears in an aperture. If these shears are coherent over the aperture, then the net result will be a large signal. The shear signals corresponding to a filament will all point in approximately the same direction along the filament's length, and thus filaments are well suited to detection using this method. However, this method is not tailored to exclusively find filaments, and a variety of mass distributions will give a strong signal. The LSS adds noise to the shear monopole signal which can, in the cases of weak filaments or low galaxy densities, cause confusion, false-identification and obscuration of the filament signal. However, we have shown that at galaxy densities characteristic of future space-based missions, or if particularly strong filaments are targeted, the filament signal will still stand out clearly against the LSS noise.  We demonstrated that one can use MCMC techniques to fit filament profiles, even at relatively low galaxy densities. This is an important result as it shows that weak lensing can be used not only for filament detection, but also for quantitative analysis of filament properties. In the future, this will allow us to directly compare the predictions of simulations to observations. The scatter in the best fit parameters caused by LSS has been quantified, and hinders the ability of this method to quantify profiles accurately - in future work, the LSS should be incorporated into the likelihood function to minimize the impact on the results (e.g. as per the prescription laid out in \\citet{HuWh}). The MCMC method was also shown to be quite versatile. It can be used on its own, as it was when we looked for filaments between cluster pairs, or it can be used in conjunction with other techniques such as the mass reconstruction or in particular the shear filter to first of all search for a filament candidate. There are a number of caveats associated with using MCMC fits. Firstly, it is model-dependent - this involves assuming a profile for the filament, in this paper based on the findings of \\cite{Colfil}. Models that more accurately reflect the `true' filament profiles will achieve better likelihood values. Secondly, since it is very difficult to write down a model for irregular or warped filaments, this method is only suitable for investigating the properties of straight, regular filaments. This represents only a fraction of the total filament population. Thirdly, the method works best if we have some prior information on the location of the filament, either from one of the other detection methods, or by using cluster pairs between which filaments are thought to be constrained. Without a prior on the filament position, fitting any profile would become a complicated procedure.  Stacking is a technique that can be used in conjunction with any of the methods described above, to improve the chance of filament detection. Stacking is particularly useful when each individual cluster lenses a low background galaxy density, but the survey contains many such clusters. In this paper we suggested two possible stacking methods: the first aligned cluster major axes in the stacked images and the second aligned axes joining cluster pairs. The latter method is far more effective. This is because, although filaments are often observed to align with the major axis of a cluster, there is no guarantee that they are either regular or strong. In almost all cases there is imperfect alignment of the filaments with the cluster major axis, and sometimes there is no filament aligned in this direction at all. If irregular or misaligned filaments are stacked the net effect will be to produce a low density mass sheet. This is not ideal for any of the methods described in this paper. On the other hand, filaments between cluster pairs tend to be straight, strong and regular and hence are easily stacked. The only disadvantage of using cluster pairs is their relative scarcity - whereas any cluster can be used when aligning the major axes, fewer cluster pairs will exist in a survey.  In summary, in this paper we have presented a variety of new detection methods, some of which (for example the shear filter and MCMC) give promising results. However, due to the low surface densities of filaments, and the additional effects of LSS noise, the task of detecting filaments in actual data remains difficult. In reality, filament detection will either require the high galaxy densities of future space-based missions, or the stacking of a number of clusters. Current space-based lensing studies achieve $\\approx$ 80 galaxies arcmin$^{-2}$ (e.g. \\citep{Leonard}), and we could achieve in excess of 250 arcmin$^{-2}$ for deep space-based observations (e.g. \\citep{Rhodes}). Although there is a trade-off between the area and depth of a survey, so that those targeting many thousands of square degrees will typically not go so deep, the most massive clusters such as those considered in this work will be prime targets for deep targeted observations. Stacking the data from the fields of less massive clusters in non-targeted large area surveys will also be extremely useful in statistically constraining the properties of filaments.  There is one further method, not discussed in this paper, that could prove valuable in detecting filaments -  gravitational flexion. In traditional weak lensing studies, such as the formalism used in this paper, the lens equation is approximated as linear. This means that any distortion in the ellipticity of a galaxy will be aligned purely tangentially to the lens. Physically, such a linear approximation is equivalent to assuming no variation in the lens field over the scale of the lensed image. However, if we do not make such an assumption, the lens equation becomes non-linear, and in addition to the convergence and shear, gravitational distortion is determined by the first and second flexion. Taken together, flexion introduces `arciness' to the image and some radial alignment with respect to the lens (see \\cite{GB05} for further details). It has been demonstrated, for example by \\cite{LKW09}, that gravitational flexion is an effective probe of substructure in galaxy clusters. This suggests that flexion may also be a useful tool for filament detection.  Besides detecting filaments using weak lensing data, as we have noted in the Introduction, there are also a variety of other complementary detection methods. Sunyayev-Zeldovich observations could in future be used in combination with more established methods of filament detection (e.g. weak lensing, X-ray, galaxy overdensity) to provide more optimal algorithms for detection and even stronger constraints on profiles. The area of overlap between DES and the South Pole Telescope, SPT, will result in several tens of thousands of galaxy clusters in common. The ability to probe the total matter content of clusters and their environments, coupled with deep optical and infrared imaging and spectroscopic observations, will enable us to determine galaxy bias as a function of environment. It will also improve our understanding of the formation of the most massive bound objects in the universe. Weak lensing of background galaxy populations becomes less effective for use as a probe of higher redshift structures \\citep{lewking} and beyond $z\\sim 1$ an alternative source population will be essential -- such as the 21cm emission from high redshift proto-structures that will be imaged with SKA, or the CMB as seen by a future mission. The ability to trace the evolution of filaments in the cosmic web will have important consequences for models of structure formation e.g. how much material is bound in filaments as a function of epoch."
    },
    "0910/0910.2883_arXiv.txt": {
        "abstract": "{Prominences are partially ionized, magnetized plasmas embedded in the solar corona. Damped oscillations and propagating waves are commonly observed. These oscillations have been interpreted in terms of magnetohydrodynamic (MHD) waves. Ion-neutral collisions and non-adiabatic effects (radiation losses and thermal conduction) have been proposed as damping mechanisms.} {We study the effect of the presence of helium on the time damping of non-adiabatic MHD waves in a plasma composed by electrons, protons, neutral hydrogen, neutral helium (\\ion{He}{i}), and singly ionized helium (\\ion{He}{ii}) in the single-fluid approximation.} {The dispersion relation of linear non-adiabatic MHD waves in a homogeneous, unbounded, and partially ion\u00a1zed prominence medium is derived. The period and the damping time of Alfv\\'en, slow, fast, and thermal waves are computed. A parametric study of the ratio of the damping time to the period with respect to the helium abundance is performed.} {The efficiency of ion-neutral collisions as well as thermal conduction is increased by the presence of helium. However, if realistic abundances of helium in prominences ($\\sim 10\\%$) are considered, this effect has a minor influence on the wave damping.} {The presence of helium can be safely neglected in studies of MHD waves in partially ionized prominence plasmas.} ",
        "introduction": "Small-amplitude oscillations and propagating waves are commonly observed in both quiescent and active region prominences/filaments. They have been interpreted in terms of magnetohydrodynamic (MHD) eigenmodes of the magnetic structure and/or propagating MHD waves. The reader is referred to some recent reviews for more information about the observational and theoretical backgrounds \\citep{oliverballester02,engvold04,ballester,banerjee,engvold08}  Prominence oscillations are known to be quickly damped, with damping times corresponding to a few oscillatory periods \\citep[this topic has been reviewed by][]{oliver, mackay}. Several damping mechanisms of MHD waves have been proposed, non-adiabatic effects and ion-neutral collisions being the more extensively investigated. In order to understand in detail these effects, they have been studied in simple configurations such as unbounded and homogeneous media. \\citet{carbonell04} investigated the time damping in a homogeneous prominence medium taking non-adiabatic effects (optically thin radiation losses and thermal conduction) into account. Later on, the spatial damping was studied by \\citet{carbonell06} and the effect of a background mass flow was analyzed by \\citet{carbonell09}. Subsequently, some works have extended these previous results by considering the presence of the coronal medium \\citep{soler07,soler08,soler09NA}. The common conclusion of these investigations is that only slow and thermal waves are efficiently damped by non-adiabatic effects, while fast waves are very slightly damped and Alfv\\'en waves are completely unaffected.  On the other hand, the influence of partial ionization on the propagation and time damping of MHD waves has been also investigated in an unbounded medium. \\citet{forteza07} followed the treatment by \\citet{brag} and derived the full set of MHD equations along with the dispersion relation of linear waves in a partially ionized, single-fluid plasma \\citep[see also][]{pinto}. The presence of electrons, protons, and neutral hydrogen atoms was taken into account, whereas helium and other species were not considered. In a subsequent work \\citep{forteza08}, they extended their previous analysis by considering radiative losses and thermal conduction by electrons and neutrals. Their main results with respect to the fully ionized case \\citep{carbonell04} were, first, that ion-neutral collisions (by means of the so-called Cowling's diffusion) can damp both Alfv\\'en and fast waves but non-adiabatic effects remain only important for the damping of slow and thermal waves, and second, that there exist critical values of the wavenumber in which the real part of the frequency vanishes, so wave propagation is not possible for larger wavenumbers. Again, applications to a more complex cylindrical geometry have been also performed \\citep{soler09IN,soler09INRA}  On the basis of these previous results, it seems clear that partial ionization plays a relevant role on wave propagation in prominences. Prominences are roughly composed by 90\\% hydrogen and 10\\% helium but, to date, all the investigations considered a pure hydrogen plasma. Therefore, the effect of the presence of helium on the propagation and damping of MHD waves is still unknown and is the motivation for the present work. Here, we consider an unbounded and homogeneous prominence medium permeated by a homogeneous magnetic field. The plasma is assumed to be partially ionized, electrons, protons, neutral hydrogen, neutral helium (\\ion{He}{i}), and singly ionized helium (\\ion{He}{ii}) being the species taken into account. Recent studies by \\citet{labrosse} indicate that for central prominence temperatures, the ratio of the number densities of \\ion{He}{ii} to \\ion{He}{i} is around 10\\%, whereas the presence of \\ion{He}{iii} is negligible. This result allows us to neglect \\ion{He}{iii} in this work. Extending the works by \\citet{forteza07,forteza08}, the derivation of the basic MHD equations for a non-adiabatic, partially ionized, single-fluid plasma has been generalized by considering now five different species, allowing us to study how the presence of neutral and singly ionized helium affects their previous results.  This paper is organized as follows. The description of the equilibrium and the basic equations are given in Sect.~\\ref{sec:equations}. The results are discussed in Sect.~\\ref{subsec:results}. Finally, Sect.~\\ref{sec:conclusions} contains the conclusion of this work. ",
        "conclusions": "\\label{sec:conclusions}  In this work, we have studied the effect of helium (\\ion{He}{i} and \\ion{He}{ii}) on the time damping of thermal and MHD waves in a partially ionized prominence plasma. This is an extension of previous investigations by \\citet{forteza07,forteza08} in which helium was not taken into account. We conclude that, although the presence of neutral helium increases the efficiency of both ion-neutral collisions and thermal conduction, its effect is not important for realistic helium abundances in prominences. In addition, due to the very small \\ion{He}{ii} abundance for central prominence temperatures, its presence is irrelevant to the wave behavior. This conclusion applies both to the free propagation case and the constrained propagation by a waveguide case. Although the role of \\ion{He}{ii} (or even \\ion{He}{iii}) could be larger for typical prominence-corona transition region temperatures, the present result allows future studies of MHD waves and oscillations in prominences to neglect the presence of helium."
    },
    "0910/0910.2551_arXiv.txt": {
        "abstract": "Empirical evidence for both stellar mass black holes ($\\mbh <10^{2}\\,\\msun$) and supermassive black holes (SMBHs, $\\mbh>10^{5}\\,\\msun$) is well established. Moreover, every galaxy with a bulge appears to host a SMBH, whose mass is correlated with the bulge mass, and even more strongly with the central stellar velocity dispersion $\\sigma_c$, the $\\msigma$ relation. On the other hand, evidence for \"intermediate-mass'' black holes (IMBHs, with masses in the range 100--$10^{5}\\,\\msun$) is relatively sparse, with only a few mass measurements reported in globular clusters (GCs), dwarf galaxies and low-mass AGNs. We explore the question of whether globular clusters extend the $\\msigma$ relationship for galaxies to lower black hole masses and find that available data for globular clusters are consistent with the extrapolation of this relationship. We use this extrapolated $\\msigma$ relationship to predict the putative black hole masses of those globular clusters where existence of central IMBH was proposed. We discuss how globular clusters can be used as a constraint on theories making specific predictions for the low-mass end of the $\\msigma$ relation. ",
        "introduction": "The empirical evidence for the ubiquity of both stellar mass black holes (black hole mass $\\mbh$ of ~\\abt~$1-15\\,\\msun$) and supermassive black holes (SMBHs, $\\mbh$~$>10^{5}~\\msun$) is well established. While it is estimated that there are about $10^7-10^9$ stellar-mass black holes in every galaxy \\citep[e.g.,][]{ShapiroTeukolsky83,BrownBethe94}, every galactic bulge appears to host a SMBH. Moreover, the kinematically determined mass of these SMBHs is correlated with the mass of the bulge, and even more strongly with the central stellar velocity dispersion $\\sigma$, the $\\msigma$ relation \\citep{FerrareseMerritt00,gebhardtetal2000c,KormendyGebhardt01,Tremaine02}. Active galactic nuclei (AGN), which have long been believed to be driven by accretion around SMBHs and whose SMBH masses have been estimated from reverberation-mapping, are also consistent with the $\\msigma$ relationship \\cite[e.g.,][and references therein]{Wandel02,Bentz09}. The corollary is that the nuclear activity is a phase (or more) in the life of (at least) every galaxy with a bulge.  The $\\msigma$ relation has been extended down to black hole masses of \\abt~$10^5\\,\\msun$ \\citep{Barth05,GreeneHo06} by searching for central BHs within very low-luminosity AGN. Their dynamical studies are an unambiguous verdict on the presence of central BHs in dwarf ellipticals and very late-type spirals (e.g., NGC4395 by \\citet{GreeneHo07} and references therein). Even lower mass black holes have been inferred from non-dynamical methods for other low-luminosity AGN \\cite[e.g.,][]{Dong07}. In spite of considerable efforts, however, evidence for black holes of still lower mass, \\viz the intermediate-mass black holes (IMBHs, of masses $10^2-10^4\\msun$), is relatively sparse. Attempts to discover IMBHs in globular clusters via their X-ray emission have been on for a long time, \\cite[e.g.,][]{BahcallOstriker75}, and although the recently discovered ultra-luminous X-ray sources (ULXs) have been attributed to IMBHs \\cite[e.g.,][]{ColbertMushotsky99}, this suggestion has also been contested \\cite[e.g.,][]{Berghea08}.  Extrapolating the $\\msigma$ correlation down the mass scale predicts that IMBHs can be found in stellar systems that have velocity dispersions of $<30$~km/sec, clearly pointing to the globular clusters. Observational evidence for such black holes are only a handful, however ({\\it cf.} references in Table~\\ref{tab:sampleclusters}). Furthermore, theoretical results on IMBHs in globular clusters remain ambiguous, at best. Some theories predict the necessity of most (if not all) Galactic globular clusters to host central black holes \\citep[e.g.,][]{MillerHamilton02}, while some argue for the impossibility of globular clusters to form and/or retain black holes in their cores \\citep[e.g.,][]{Favata04,KawakatuUmemura05}. The importance of investigating the possibility of globular clusters hosting IMBHs cannot be overestimated. While a central IMBH would clearly impact the evolution of the globular cluster itself, more generally, IMBHs are crucial to link the formation processeses of stellar-mass BHs and SMBHs, and could have served as seeds for the growth of SMBHs. Extending the local BH mass function to the extreme low end can help in understanding whether there is a minimum galaxy mass or velocity dispersion below which BHs are unable to form or grow \\citep[e.g.,][]{Bromley04}. Globular clusters and galactic bulges are both `hot' stellar systems, and, since all large bulge systems seem to have a central black hole, to the extent that a massive, bound globular cluster can be viewed as a ``mini-bulge\", it may be that every dense stellar system (small or large) hosts a central black hole \\citep{Gebhardt02}. BHs at the low end of the mass ladder can be used as a constraint on theoretical models with different predictions on the $\\mbh-\\s$ behaviour. For example, the prediction that $\\msigma$ relation shall substantially steepen below $\\s=150$ km/sec \\citep{Granato04} is not supported by the low-mass AGN sample \\citep{Barth05}. The question of IMBHs also has bearing on debates in cosmology, since the cosmic mass-density of IMBHs could exceed that of SMBHs ($\\O\\approx 10^{-5.7}$), and may even account for all the baryonic dark matter in the universe, with $\\O\\approx 10^{-1.7}$ \\citep{vanderMarel03}.  The few recent reports on the detection of central black holes in some globular clusters seem to suggest that globular clusters do follow the $\\msigma$ relation for SMBHs \\citep{vanderMarel02,McLaughlin06, Noyola08,Lanzoni07,Gebhardt02} and that these BHs represent the `true' IMBHs in the mass range of $10^2-10^4\\,\\msun$. In this paper, we compile published discoveries of central black holes in globular clusters and investigate the consistency of the data with the $\\msigma$ and $\\mbh-$luminosity relations. Using these globular clusters as a constraint on the slope of the extended $\\msigma$ relation, we estimate the black hole masses for a sample of globular clusters that are proposed to host IMBHs and discuss the implications. ",
        "conclusions": "We extend the $\\msigma$ relationship for galaxies to lower black hole masses using black holes discovered in globular clusters. The reported masses of black holes in the centres of globular clusters M15, 47~Tuc, $\\omega$~Cen, NGC~6388 and G1 are consistent with the linear extrapolation of the $\\msigma$ relationship for galaxies to the low-mass end. Using this extrapolated relationship, we have estimated the masses of putative black holes in a sample of globular clusters, and find that these clusters are consistent with the $\\mbh-$luminosity relationship as well.  In the extended $\\msigma$ plot, points corresponding to different types of stellar systems occupy distinct regions, suggesting that black hole masses of \\abt~few~$\\times~10^{4}~\\msun$ represent the {\\it lowest} limit for the central black holes  of galaxies.  Masses of the central black holes that are below this limit correspond to the globular cluster domain (keeping in mind that G1 and $\\omega$~Cen are believed to be tidally stripped dwarf galaxies and not genuine globular clusters).  The consistency of black hole masses in globular clusters with the extrapolated $\\msigma$ and $\\mbh-$luminosity relationships reinforces the idea that globular clusters harbour IMBHs in their centres. The central black holes of globular clusters place even stronger constraints than low-luminosity AGNs on theoretical models of supermassive black hole growth and evolution. If black holes in globular clusters do exist, they will rule out models that predict either steepening or flattening of the $\\msigma$ relation the low-mass region.    \\appendix"
    },
    "0910/0910.0248_arXiv.txt": {
        "abstract": "\\begin{list}{ } {\\rightmargin 1in} \\baselineskip = 11pt \\parindent=1pc {\\small Direct measurements of the spectra of extrasolar giant planets are the keys to determining their physical and chemical nature.  The goal of theory is to provide the tools and context with which such data are understood.  It is only by putting spectral observations through the sieve of theory that the promise of exoplanet research can be realized.  With the new {\\em Spitzer} and HST data of transiting ``hot Jupiters,\" we have now dramatically entered the era of remote sensing. We are probing their atmospheric compositions and temperature profiles, are constraining their atmospheric dynamics, and are investigating their phase light curves.  Soon, many non-transiting exoplanets with wide separations (analogs of Jupiter) will be imaged and their light curves and spectra measured. In this paper, we present the basic physics, chemistry, and spectroscopy necessary to model the current direct detections and to develop the more sophisticated theories for both close-in and wide-separation extrasolar giant planets that will be needed in the years to come as exoplanet research accelerates into its future. \\\\~\\\\~\\\\~}%  \\end{list} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.2284_arXiv.txt": {
        "abstract": "\\fontsize{10}{10.6}\\selectfont  The Serpens SMM 1 region was observed in the 6.9 mm continuum with an angular resolution of about 0\\farcs6. Two sources were found to have steep positive spectra suggesting emission from dust. The stronger one, SMM 1a, is the driving source of the bipolar jet known previously, and the mass of the dense molecular gas traced by the millimeter continuum is about 8 $M_\\odot$. The newly found source, SMM 1b, positionally coincides with the brightest mid-IR source in this region, which implies that SMM 1b is yet another young stellar object. SMM 1b seems to be less deeply embedded than SMM 1a. SMM 1 is probably a protobinary system with a projected separation of 500 AU. ",
        "introduction": "The Serpens dark cloud is a nearby star-forming region (see Harvey et al. 2007 and references therein). The cloud core contains a cluster of low-mass young stellar objects, and some of them are extremely young protostars (Testi \\& Sargent 1998; Harvey et al. 2007; Winston et al. 2007). The core also harbors several Herbig-Haro objects and molecular outflows (Davis et al. 1999; Hodapp 1999). The most luminous and most deeply embedded object among them is SMM 1 (Harvey et al. 1984; Casali et al. 1993; Enoch et al. 2007). SMM 1 is a Class 0 source and associated with a bipolar jet/outflow (Rodr{\\'\\i}guez et al. 1989; Curiel et al. 1996; White et al. 1995; Larsson et al. 2000).  Because of the extremely high extinction, the small scale structure of SMM 1 has been accessible only with radio interferometric observations. Images of centimeter continuum revealed a well-collimated bipolar jet showing a large proper motion away from the central protostar, indicating a dynamical age of $\\sim$60 yr (Rodr{\\'\\i}guez et al. 1989; Curiel et al. 1993). Hogerheijde et al. (1999) imaged SMM 1 in the millimeter continuum and found a single compact source surrounded by an envelope of $\\sim$9 $M_\\odot$. Hogerheijde et al. (1999) also found that the molecular outflow driven by SMM 1 has a complicated structure.  One of the interesting issues in the study of star formation is the multiplicity, because the majority of stars belong to multiple-star systems (Duquennoy \\& Mayor 1991). The multiplicity of young pre-main-sequence stars is higher than that of main-sequence stars (Duch{\\^e}ne et al. 2007), which implies that multiple systems form in the early stage of star formation. In fact, bright protostars, such as IRAS 16293--2422 and NGC 1333 IRAS 4A, are usually found to be multiple systems when imaged with a high angular resolution (Looney et al. 2000). Serpens SMM 1 has been considered as an interesting example of an isolated protostar that can be understood well because it is relatively bright, massive, and simple. However, there were some indications of complexity, such as the near-IR excess (Larsson et al. 2000). Therefore, it is necessary to investigate the structure of SMM 1 with a high angular resolution and a high sensitivity.  In this paper, we present the results of our observations of the Serpens SMM 1 region in the 6.9 mm continuum with the Very Large Array (VLA) of the National Radio Astronomy Observatory. We describe our radio continuum observations and archival data in Section 2. In Section 3 we report the results of the continuum imaging. In Section 4 we discuss the star-forming activities in the SMM 1 region. ",
        "conclusions": ""
    },
    "0910/0910.4970_arXiv.txt": {
        "abstract": "We report on the measurement of the physical properties (rest frame K-band luminosity and total stellar mass) of the hosts of 89 broad line (type--1) Active Galactic Nuclei (AGN)  detected in the zCOSMOS survey in the redshift range $1<z<2.2$. The unprecedented multi-wavelength coverage of the survey field allows us to disentangle the emission of the host galaxy from that of the nuclear black hole in their Spectral Energy Distributions (SED). We derive an estimate of black hole masses through the analysis of the broad \\MgII\\ emission lines observed in the medium-resolution spectra taken with {\\it VIMOS/VLT} as part of the zCOSMOS project. We found that, as compared to the local value, the average black hole to host galaxy mass ratio appears to evolve positively with redshift, with a best fit evolution of the form $(1+z)^{0.68 \\pm 0.12 ^{+0.6}_{-0.3}}$, where the large asymmetric systematic errors stem from the uncertainties in the choice of IMF, in the calibration of the virial relation used to estimate BH masses and in the mean QSO SED adopted. On the other hand, if we consider the observed rest frame K-band luminosity, objects tend to be brighter, for a given black hole mass, than those on the local $M_{\\rm BH}$-$M_K$ relation. This fact, together with more indirect evidence from the SED fitting itself, suggests that the AGN hosts are likely actively star forming galaxies. A thorough analysis of observational biases induced by intrinsic scatter in the scaling relations reinforces the conclusion that an evolution of the $M_{\\rm BH}-M_{*}$ relation must ensue for actively growing black holes at early times: either its overall normalization, or its intrinsic scatter (or both) appear to increase with redshift. This can be interpreted as signature of either a more rapid growth of supermassive black holes at high redshift, a change of structural properties of AGN hosts at earlier times, or a significant mismatch between the typical growth times of nuclear black holes and host galaxies. In any case, our results provide important clues on the nature of the early co-evolution of black holes and galaxies and challenging tests for models of AGN feedback and self-regulated growth of structures. ",
        "introduction": "Tight scaling relations between the central black holes mass and various properties of their host spheroids (velocity dispersion, $\\sigma_{*}$, stellar mass, $M_{*}$, luminosity, core mass deficit) characterize the structure of nearby inactive galaxies \\citep{magorrian:98,gebhardt:00,ferrarese:00,tremaine:02,marconi:03,haering:04,hopkins:07b,kormendy:09,gultekin:09}. A result of the search for local QSO relics via the study of their dynamical influence on the surrounding stars and gas made possible by the launch of the {\\it Hubble Space Telescope} nearly twenty years ago, these correlations have revolutionized the way we conceive the physical link between galaxy and AGN evolution. Coupled with the fact that supermassive black holes (SMBH) growth is now known to be due mainly to radiatively efficient accretion over cosmological times, taking place during ``active'' phases \\citep{soltan:82,marconi:04,shankar:04,merloni:08}, this led to the suggestion that most, if not all, galaxies went through a phase of nuclear activity in the past, during which a strong physical coupling (generally termed 'feedback') must have established a long-lasting link between host's and black hole's properties.  This shift of paradigm has sparked the activity of theoretical modelers. Following  pioneering analytic works (Ciotti \\& Ostriker 1997,2001; Silk \\& Rees 1998; Fabian 1999; Cavaliere \\& Vittorini 2002; Whyithe \\& Loeb 2003), widely different approaches have been taken to study the role of AGN in galaxy evolution. Semi-analytic models (SAM) have been the most numerous (Monaco et al. 2000; Kauffmann \\& Haehnelt 2000; Volonteri et al. 2003; Granato et al. 2004; Menci et al. 2006; Croton et al. 2006; Bower et al. 2006; Malbon et al. 2007; Marulli et al. 2008; Somerville et al. 2008), and, among other things, have helped establishing the importance of late-time feedback from radio-active AGN for the high-mass end of the galaxy mass function. However, firmer conclusions on the physical nature of such a feedback mode have been hampered by the large freedom SAM have in choosing baryonic physics recipes to implement in their schemes. On the other hand, fully hydrodynamic simulations of the cosmological evolution of SMBH have been also performed (Di Matteo et al. 2003, 2008; Sijacki et al. 2007; Coldberg \\& Di Matteo 2008), but their computational costs have so far allowed only a limited exploration of sub-grid prescriptions (AGN physical models) in relatively small cosmological volumes. A third, hybrid, approach has also been followed, in which the results of high-resolution hydrodynamic simulations of galaxy-galaxy mergers with black holes (Di Matteo et al. 2005; Springel et al. 2005) have been used to construct a general framework for merger-induced AGN feedback, capable of passing numerous observational tests (Hopkins et al. 2006).  Almost all the approaches outlined above use the local scaling relations as a constraint to the model parameters. As such, these relations have proved themselves unable to unambiguously determine the physical nature of the SMBH-galaxy coupling. One obvious way out of this impasse is to study their redshift evolution, that different models predict to be different \\citep{granato:04,robertson:06,croton:06b,fontanot:06,malbon:07,marulli:08,hopkins:09a}.  From the observational point of view, the current situation is far from being clear. In recent years, a number of groups have employed different techniques to try and detect signs of evolution in any of the locally observed scaling relations. Most efforts have been devoted to the study of the $M_{\\rm BH}-\\sigma_{*}$ relation. Shields et al. (2003) and Salviander et al. (2007) have used narrow nebular emission lines ([OIII], [OII]) excited by the AGN emission in the nuclear region of galaxies as proxies for the central velocity dispersion, and compared these to the black hole mass estimated from the broad line width of QSOs from $z\\sim 0$ to $z\\sim 3$ (see \\S \\ref{sec:virial}). In both cases, a large scatter has been found in the relation between $M_{\\rm BH}$ and $\\sigma_{*}$, and the results  are either in favor of (Salviander et al. 2007) or against \\citep{shields:03} a positive evolution of the black hole mass to host dispersion ratio. However, as pointed out by Botte et al. (2004) and Greene \\&  Ho (2005), there are a number of problems with the underlying assumption that the narrow emission lines are good probes of the central gravitational potential, and the systematic uncertainties this method is endowed with are large. Komossa and Xu (2007) have also shown that the the AGN for which the [OIII] emission line width is far broader that the host galaxy's stellar velocity dispersion $\\sigma_{*}$ tend to show clear sign of blue-shift in the narrow emission line (of the order of 100 km/s or larger). Thus, using [OIII] as a proxy for $\\sigma_{*}$ would at least require good enough spectral resolution to measure such blue-shifts.  An alternative path has been followed by Woo et al. (2006), Treu et al. (2007) and Woo et al. (2008), who have studied carefully selected samples of moderately bright AGN in narrow redshift ranges ($z\\sim 0.36$ and $0.57$), where the host's stellar velocity dispersion can be measured directly from the absorption lines in high signal-to-noise spectra. They also found evidence of (strong) positive evolution of the $M_{\\rm BH}$ to $\\sigma_{*}$ ratio compared to the local value. This method, although promising and reliable, is quite inefficient and telescope-time consuming: secure detection of spectral absorption features in massive ellipticals at $1\\la z \\la 2$ require hundreds of hours of integration time on a 8-meter class telescope \\citep{cimatti:08}.  Other groups have chosen to try and derive information on the host mass of broad line AGN using multi-colour image decomposition techniques \\citep{jahnke:04,sanchez:04} or spatially deconvolving optical spectra \\citep{letawe:07}. Due to the severe surface brightness dimming effects, employing these techniques for redshift QSOs becomes increasingly challenging (but see Schramm, Wisotzki and Jahnke 2008; Jahnke et al. 2009; Decarli et al. 2009, in prep.), unless gravitationally lensed QSOs are selected \\citep{peng:06b,ross:09}. In all cases, very deep, high resolution optical images ({\\it HST}) are necessary to reliably disentangle the nuclear from the host galaxy emission. Statistically, the most significant results have been published by Peng et al. (2006b), who, based on a large sample of 51 AGN (both lensed and nonlensed) in the range $1<z<4.5$, observed that the ratio $M_{\\rm BH}/M_{*}$ increases with look-back time, up to a factor $\\simeq 4^{+2}_{-1}$ at the highest redshift probed.  Also cold/molecular gas motions on large galactic scales have been used to infer the properties (total mass in particular) of the hosts of AGN at different redshifts. In the local Universe, HI 21 cm lines have been used by Ho et al. (2008) to derive total stellar masses of the bulges around Seyfert-like AGN. Using instead CO lines to estimate the velocity dispersion in high-redshift QSO hosts, Walter et al. (2004) and Shields et al. (2006) find tentative evidence that the hosts are very under-massive compared to their central BHs (i.e., a positive redshift evolution of the  $M_{\\rm BH}/M_{*}$ ratio), a result confirmed by the study of Ho (2007).  Finally, a completely different approach has been that of trying to follow and compare the evolution of global descriptors of the galaxy and SMBH populations, such as mass densities (Merloni, Rudnick and Di Matteo 2004; Hopkins et al. 2006b; Shankar et al. 2008). Using the simple {\\it ansatz} that total black hole mass density can only increase with time, and requiring that the limits imposed by local demographics are not violated, these works showed in general very moderate, if any, signs of cosmological evolution of the average black hole to host mass ratio.  Here we present a new method to tackle the issue of studying black hole-galaxy scaling relations at high redshift. Starting from a sample of un-obscured AGN, for which the broad line kinematics can be used to infer the central SMBH mass, we take advantage of the unprecedentedly deep multi-wavelength coverage of the COSMOS \\citep{scoville:07} field and develop a novel SED fitting technique that allows us to decompose the entire spectral energy distribution into a nuclear AGN and a host galaxy components \\citep{bongiorno:07}. We show here how, for the majority of the objects in our sample, rest frame K-band luminosity and total stellar mass of the host can be robustly determined, opening the way to a detailed study of the scaling relations in type--1 AGN at $1\\la z \\la 2.2$.  The structure of the paper is the following: we will begin (\\S~\\ref{sec:sample}) by introducing our sample, before proceeding to a discussion of our SED fitting method in section~\\ref{sec:sed_fit}, focusing our attention on the measures of the rest-frame K-band luminosity of the AGN hosts (\\S~\\ref{sec:mk}), on their total stellar mass (\\S~\\ref{sec:mstar}). In section~\\ref{sec:virial} we will describe our estimates of the black hole masses obtained by studying the properties of the broad \\MgII\\ emission line, while in \\S~\\ref{sec:bhmass} we will briefly outline the characteristics of our AGN sample in terms of bolometric luminosity, BH mass and accretion rates. Section~\\ref{sec:scaling_evol} contains the main novel results of our study, namely the analysis of the scaling relation for the objects in our sample, as well as a characterization of their observed redshift evolution. An important part of our analysis, however, is the assessment of possible observational biases responsible for the trend observed, that we carry out in section~\\ref{sec:bias}. This allows us to reach robust conclusions, that we discuss at the end of the paper, in section~\\ref{sec:disc}.  Throughout this paper, we use the standard cosmology ($\\Omega_m=0.3$, $\\Omega_{\\Lambda}=0.7$, with $H_0=$70 km s$^{-1}$ Mpc$^{-1}$). ",
        "conclusions": "\\label{sec:disc}  \\subsection{Our results} \\label{sec:our_results} We have used an AGN+host galaxy SED decomposition technique to infer the physical properties of the hosts of 89 (moderately luminous: i.e. mostly sub-$L*$) type--1 AGN in the zCOSMOS survey. Thanks to the deep, intensive multi-wavelength coverage of the COSMOS field, the observed spectral energy distributions are sampled to such a degree that the decomposition technique works reasonably well. We are thus able to derive rest frame K-band magnitudes and total stellar masses for the majority of our sample (80-90\\%, depending on the choice of AGN SED). Noticeably, our method allows us to properly quantify the uncertainty in each measurement.  The bulk of the sample of BLAGN hosts we have studied have total stellar masses in the range 10$^{10.5}$--10$^{11.3} M_{\\odot}$. Both the derived mass-to-light ratios and the sample average star formation rates seem to suggest that they are moderately-to-highly star forming objects, in good agreement with the host properties of both type 1 and type 2 (obscured) X-ray and/or IR selected AGN, either in a lower (Kauffmann et al. 2003; Jahnke et al. 2004; Hickox et al. 2009; Silverman et al. 2009) and in a similar (Brusa et al. 2009; Ross et al. 2009) redshift range. Reassuringly, also the fits with the P07 phenomenological templates indicate an overall preference for star-forming hosts over passive (elliptical) ones. Obviously, a more accurate determination of the individual star formation rates is hampered by the dominant contribution of the AGN in the rest-frame UV/Optical part of the spectra.  The black hole masses derived with the the ``virial'' or ``empirically calibrated photo-ionization'' method are broadly distributed around Log$M_{\\rm BH}\\sim 8.5$. Interestingly, the mass range probed by our sample corresponds to that where the most stringent constraints on the $z=0$ scaling relations are available. The Eddington ratios $\\lambda$ of our objects have a mean of Log$\\lambda\\approx -1$. We observe a clear trend between Eddington ratio and bolometric luminosity, that could be indicative of some specific luminosity-dependent AGN lifetimes distribution \\citep{hopkins:09b}, but the underlying effect of various selection biases in determining this trend needs to be further assessed \\citep{trump:09}.  For a large number of objects in the sample (68/89; but the number falls to 24/89 if we consider the errors) the measured black hole to stellar host mass ratio is positively offset from that predicted by the local H{\\\"a}ring \\& Rix (2004) local scaling relation. Assuming a redshift-dependent evolution in the $\\Delta{\\rm Log}(M_{\\rm BH}/M_{*})$ of the form $\\delta_2{\\rm Log}(1+z)$, we measure $\\delta_2=0.68 \\pm 0.12 ^{+0.6}_{-0.3}$, where the large asymmetric systematic errors stem from the uncertainties in the hosts' IMF, in the calibration of the virial relation used to estimate BH masses and in the mean QSO SED to be used.  The scatter in the measured offset is substantial at all redshifts probed, such that we could still be consistent with a lack of evolution in the scaling relation at the 2-$\\sigma$ level, as shown by the inset of figure~\\ref{fig:mbh_mk_mass}. There, it is also apparent that the majority of observational data-points for AGN samples in the published literature do indeed show a broadly consistent amount of offset at all redshifts probed. This might suggest that intrinsic differences in the SMBH/host galaxy relation between active and inactive galaxies could play an important role besides any genuine cosmological evolution.  Jahnke et al. (2009) have independently computed total stellar masses for a small sample (18 objects) of X-ray selected AGN in the COSMOS field for which simultaneous {\\it HST}/ACS and {\\it HST}/NICMOS observations allow to clearly image the resolved hosts. The two independent methods for measuring stellar masses are in broad agreement with each other, with a dispersion well within the (large) uncertainties that characterize each of these methods. 5/18 of them have also zCOSMOS spectra, and are part of the sample described here (empty stars in Fig.~\\ref{fig:mbh_mk_mass}). However, the objects in the Jahnke et al. (2009) sample do not show any significant offset from the local scaling relation. This might be due to a statistical fluctuation (more so given that they tend to lie in the lower redshift range probed by our sample, where the offset from the local scaling relation is smaller), or may indicate a more serious issue with the different estimates of the mass-to-light ratio of the AGN hosts. Larger samples of {\\it HST}-imaged AGN hosts at longer wavelengths (with the newly installed WFC3) will be extremely important to settle this issue and, in general, to improve the calibration of our SED-based method to estimated host galaxy masses for AGN.  We have taken particular care in examining the effects of the bias inevitably introduced by any intrinsic scatter in the BH-host mass scaling relation into any AGN-selected sample, as ours. We conclude that our data cannot possibly be explained if type--1 AGN and their hosts at $1<z<2$ lie on a scaling relation which has the same slope, normalization and scatter as the locally observed one. On the other hand positive evolution of the average $M_{\\rm BH}/M_{*}$ ratio (larger black holes at early times in unobscured AGN) or of the intrinsic scatter (or a combination of the two) are needed to explain our results.   \\subsection{Implications for theoretical models} What are the implications of these findings for our understanding of the cosmological co-evolution of black holes and galaxies? Let us briefly discuss recent theoretical investigations on this issue and the corresponding predictions for the evolution of the scaling relations.  One of the earliest semi-analytic models to incorporate the evolution of supermassive black holes and the associated feedback effect were those by Granato et al. (2001,2004). There, triggering of AGN activity is not directly linked to merger activity, but rather generically to the process of bulge/spheroid formation. The rate of star formation and black hole accretion is regulated by the starlight radiation drag, and consequently, in a typical system the ratio $M_{\\rm BH}/M_{*}$ is initially small and rapidly grows until the AGN feedback sweeps the remaining gas. QSOs and, in general, type--1 AGN are thus associated to the final stage of bulge formation, and it is very hard to produce any positive offset from the local relation like the one we measure.  This is however a problem common to all feedback models in which the black hole energy injection is very fast (explosive). Indeed, the first published predictions of merger-induced AGN activity models \\citep{robertson:06} indicated that, if strong QSO feedback is responsible for rapidly terminating star formation in the bulge \\citep{dimatteo:05,springel:05}, then very little evolution, as well as very little scatter, is expected for the scaling relations. However, later works within the same theoretical framework \\citep{hopkins:07b,hopkins:09a} have analyzed in greater depths the role of dissipation in major mergers at different redshift. Under the assumption that black hole and spheroids obey a universal ``black hole fundamental plane'' (BHFP), where $M_{\\rm BH}\\propto M_{*}\\sigma_{*}^2$, they show how, in gas-richer environments (at higher redshift), dissipation effects may deepen the potential well around the black hole, allowing it to grow above the $z=0$ $M_{\\rm BH}-M_{*}$ relation, to a degree marginally consistent with our results. However, in the same physical framework, the $M_{\\rm BH}-\\sigma_{*}$ relation is almost independent on redshift, which would contradict the observational results of \\cite{woo:06,treu:07,woo:08}.  Another, related effect was discussed in Croton (2006). There it was assumed that major mergers can trigger both star formation in a bulge as well as black hole growth, in a fixed proportion. However, bulges can also acquire mass by disrupting stellar discs, a channel that should not contribute to black hole growth. The relative importance of these two paths of bulge formation may lead to lighter bulges for a given black hole mass at high redshift, as disks have a smaller stellar fraction. A subsequent study of this and other dynamical process of disk-to-bulge transformation was included in the work by Fontanot et al.~(2006) and Malbon et al.~(2007). They also confirmed qualitatively the predictions of Croton (2006), but found a much smaller effect, at most a factor $\\sim 2$ at $z=2$ in the Malbon et al. (2007) work, and preferentially for small mass black holes $M_{\\rm BH} \\la 10^{8} M_{\\odot}$. Even more complex semi-analytic models including various flavors of AGN-driven winds and their feedback effects \\citep{fontanot:06} can lead to various degrees of positive redshift evolution of the average $M_{\\rm BH}/M_{*}$ ratio (Lamastra et al. 2009).  As a general rule, we observe that, following increasing observational efforts to study the evolution of scaling relations, semi-analytic models have become more sophisticated over the years. This increase of sophistication's has allowed more complex behaviours of the coupled black holes-galaxy systems over cosmological times. As it is expected, more complex models also lead to an increase in the predicted scatter, even though a clear theoretical study on the redshift evolution of such scatter is still missing (but see the recent attempts by Lamastra et al. 2009 and Somerville 2009).  Hints from hydrodynamical simulations, both of isolated mergers \\citep{johansson:09} and of relatively small cosmological boxes \\citep{coldberg:08} do indeed show a large scatter in the instantaneous ratio between black hole accretion and star formation rates, similar to what was found here (see also Silverman et al. 2009), thus suggesting that on the relatively short timescales over which un-absorbed AGN/QSOs are visible the physical connection between black holes growth and galaxy formation must be complex, too. How this would impact on the statistical and evolutionary properties of the galaxy population as a whole, however, is far from clear.  The results we have presented in \\S \\ref{sec:scaling_evol}, coupled with analysis of selection biases of \\S \\ref{sec:sel_bias} would suggest that a greater effort should be made by theoretical modellers to include a more accurate and realistic study of the evolution of the intrinsic scatter in any scaling relations, as well as that of slope and normalization.  The ``increased scatter'' hypothesis, as an explanation of the observed offset, and its inevitable consequence that a 'luminosity function weighted' bias plays a significant role in the observed evolution of scaling relations, could be strengthened if recent claims of under-massive black holes in IR selected galaxy samples were confirmed \\citep{shapiro:09}. Indeed, we should expect that in samples selected purely on the basis of the host galaxy stellar mass, rather than on AGN properties, the intrinsic scatter in the scaling relation should produce a bias going in the opposite direction as those discussed above, depending on the exact shape of the BH mass function at the redshift considered.   \\subsection{Concluding remarks} By taking advantage of the unique combination of VLT spectroscopy and deep multi-wavelength coverage of the COSMOS field, we have presented here a novel method to study the physical link between supermassive black holes and their host galaxies in type--1 (un-obscured) AGN in the crucial redshift range $1\\la z \\la 2$. The main focus of this work is on the capability of our SED decomposition technique to provide reliable estimates of the total stellar mass of the AGN hosts, and, even more importantly, reliable estimates of its uncertainty.  The main result of our study is the observation of an offset in the $M_{\\rm BH}-M_{*}$ relation, such that, in the redshift range probed, for their given hosts black holes are on average 2-3 times larger than their counterparts in the nuclei of nearby inactive galaxies. A thorough analysis of all possible observational biases induced by intrinsic scatter in the scaling relations reinforces the conclusion that an evolution of the $M_{\\rm BH}-M_{*}$ relation must ensue for actively growing black holes at early times: either its overall normalization, or its intrinsic scatter (or both) must increase significantly with redshift.  We close with two recommendations for future studies of the subject. From the observational point of view, it will be very important to explore methods to derive robust black hole mass estimates in high redshift samples of obscured AGN, that can be selected purely on the basis of their host galaxy properties. Broad emission lines at longer wavelengths, where the effect of obscuration are less severe, could be very useful in this respect. Also, a better understanding of the differences in the hosts' properties of active and inactive black holes is needed to allow a more meaningful comparison with the local scaling relations, and a better assessment of their evolution. From the theoretical point of view, more efforts should be devoted to derive robust predictions for the coupled evolution of slope, normalization {\\it and intrinsic scatter} in the scaling relations, and to properly include in the models the selection effects that clearly play an often decisive role in the observational studies of co-evolving galaxies and black holes."
    },
    "0910/0910.3885_arXiv.txt": {
        "abstract": "We have obtained deep infrared $J$ and $K$ band observations of nine $4.9 \\times 4.9$ arcmin fields in the Small Magellanic Cloud (SMC) with the ESO New Technology Telescope equipped with the SOFI infrared camera. In these fields, 34 RR Lyrae stars catalogued by the OGLE collaboration were identified. Using different theoretical and empirical calibrations of the infrared period-luminosity-metallicity relation, we find consistent SMC distance moduli, and find a best true distance modulus to the SMC of $18.97 \\pm 0.03$ (statistical) $\\pm$ 0.12 (systematic) mag which agrees well with most independent distance determinations to this galaxy, and puts the SMC 0.39 mag more distant than the LMC for which our group has recently derived, from the same technique, a distance of 18.58 mag. ",
        "introduction": "In our ongoing Araucaria Project (e.g. Gieren et al. 2005a), we are applying a number of different stellar standard candles to independently determine the distances to a sample of nearby galaxies.  The systematic differences between the distance results obtained for the individual galaxies from the various stellar candles will be analyzed in forthcoming papers. This analysis is expected to lead to a detailed understanding of how the various stellar techniques, which are fundamental to calibrate the first rungs of the distance ladder, depend on metallicity and age.  While the objects we use for the distance determinations are usually detected from optical wide-field imaging surveys of the target galaxies (e.g. Pietrzy{\\'n}ski et al. 2002b), the most accurate distance work is then done from follow-up near-infrared images which virtually eliminate reddening as a significant source of error on the results. Examples of this very successful approach are the Cepheid work on the Sculptor galaxies NGC 300 (Gieren et al. 2005b), NGC 55 (Gieren et al. 2008) and NGC 247 (Gieren et al. 2009), and the red clump star work on the LMC (Pietrzy{\\'n}ski \\& Gieren 2002). The Araucaria Project has also been developing a new spectroscopic distance indicator, viz. the Flux-weighted Gravity- Luminosity Relationship (FGLR) for blue supergiants (Kudritzki et al. 2003; 2008) which is able to yield distances accurate to 5\\% to galaxies containing massive blue stars out to about 10 Mpc from low-resolution spectra (Kudritzki et al. 2008, Urbaneja et al. 2008).  Thanks to several recent theoretical and empirical studies, evidence has been mounting that RR Lyrae (RRL) stars are excellent standard candles in the near-infrared spectral range, providing distance results which are superior to the traditional optical method (e.g. Bono 2003a).  Longmore et al. (1986) were the first to show that RRL variable stars follow a period-luminosity (PL) relation in the near-infrared K-band.  Their pioneering work was followed by Liu \\& Janes (1990), Jones et al. (1996), and Skillen et al. (1993), who applied infrared versions of the Baade-Wesselink method to calibrate the luminosities and distances of RRL stars.  A very comprehensive analysis of the IR properties of RRL stars was given by Nemec et al. (1994).  The first theoretical constraints on the K-band PL relation of RRL stars are based on non-linear convective pulsation models which were presented by Bono et al. (2001).  Dall'Ora et al. (2004) later demonstrated that the K-band PL relation for RRL stars appears to have a very small scatter for globular clusters, with small intrinsic spread in metallicity for this type of stars. Further theoretical explorations of the RRL period-mean magnitude-metallicity relations in near-infrared passbands were carried out by Bono et al. (2003b), Catelan et al. (2004), and Cassisi et al. (2004). Most recently, Sollima et al. (2008) analyzed near-infrared K-band data of RRL stars in some 15 Galactic globular clusters and provided the first empirical calibration of the period-luminosity-metallicity (PLZ) relation in the $K$ band. All these existing theoretical and empirical studies have suggested that the RRL star K-band PLZ relation appears to be a superb means to determine accurate distances to galaxies hosting an abundant old stellar population.  We have therefore started to include this tool in the Araucaria Project distance work.  In previous papers, we determined the distance to the Sculptor dwarf galaxy (Pietrzy{\\'n}ski et al. 2008), and to the LMC (Szewczyk et al. 2008) from this method. In this paper, we apply the technique to a sample of RR Lyrae stars distributed over the SMC, and determine the distance to the SMC for the first time with the infrared RR Lyrae technique. ",
        "conclusions": "The results of our deep infrared imaging and photometry of 34 RR Lyrae stars spread over nine $4.9 \\times 4.9$ arcmin fields in the SMC are presented.  Our data show two clear sequences in the period luminosity plane, which correspond to the RRc and RRab stars. After fundamentalizing the RRc periods to the period of the RRab stars by adding $\\log P = 0.127$, both groups were merged, and the distance moduli to the SMC were determined using different theoretical and empirical calibrations.  Our final adopted distance agrees very well with modern results obtained from a number of different techniques.  Our results confirm that the RR Lyrae period-luminosity-metallicity relations in the near-infrared passbands are potentially a very good tool for accurate distance measurements. In future programs, particularly for the LMC, we want to reduce the systematic uncertainty of our distance determinations with this method by providing multi-phase JK observations of the RR Lyrae variables to determine their mean magnitudes with better precision."
    },
    "0910/0910.4141_arXiv.txt": {
        "abstract": "{ The recent observational evidence for the current cosmic acceleration have stimulated renewed interest in alternative cosmologies, such as scenarios with interaction in the dark sector (dark matter and dark energy). In general, such models contain an unknown negative-pressure dark component coupled with the pressureless dark matter and/or with the baryons that results in an evolution for the Universe rather different from the one predicted by the standard $\\Lambda$CDM model. In this work we test the observational viability of such scenarios by using the most recent galaxy cluster gas mass fraction versus redshift data ($42$ X-ray luminous, dynamically relaxed galaxy clusters spanning the redshift range $0.063 < z < 1.063$) \\citep{Allen08} to place bounds on the parameter $\\epsilon$ that characterizes the dark matter/dark energy coupling. The resulting are consistent with, and typically as constraining as, those derived from other cosmological data. Although a time-independent cosmological constant ($\\Lambda$CDM model) is a good fit to these galaxy cluster data, an interacting dark energy component cannot yet be ruled out. ",
        "introduction": "Although fundamental to our understanding of the Universe, several important questions involving the nature of the dark matter and dark energy components and their roles in the dynamics of the Universe remain unanswered. Among these questions, the possibility of interaction in the dark sector, which gave origin to the so-called models of coupled quintessence, has been largely explored in the literature. These scenarios are based on the premise that, unless some special and unknown symmetry in Nature prevents or suppresses a non-minimal coupling between these components (which has not been found), such interaction is in principle possible and, although no observational piece of evidence has so far been unambiguously presented, a weak coupling still below detection cannot be completely excluded (see \\citealt{amendola};\\citealt{zimd,barrow,wu,ernandes,ernandes1}).  From the observational viewpoint, these models are capable of explaining the current cosmic acceleration, as well as other recent observational results. From the theoretical point of view, however, critiques to these scenarios do exist and are mainly related to the fact that in order to establish a model and study their observational and theoretical predictions, one needs first to specify a phenomenological coupling between the cosmic components.  In this work, instead of adopting the traditional approach, we follow the qualitative arguments used in {\\citep{Wang05, Alcaniz05}} and deduce the interacting law from a simple argument about the effect of the dark energy decay on the dark matter (CDM) expansion rate. The resulting expression is a very general law that has many of the previous phenomenological approaches as a particular case~(see, e.g., {{\\citealt{Wang05}} for a discussion). ",
        "conclusions": "We have used gas mass fraction data to test the viability of a general class of coupled quintessence models. In particular, we have proposed a new low-$z$ test which in principle may be able to distinguish classes of coupled and uncoupled dark matter/energy models. We believe that with new and more precise data sets we will be able to use this test to place tighter and stronger bounds on the interacting parameter $\\epsilon$."
    },
    "0910/0910.4007_arXiv.txt": {
        "abstract": "{% We present the Fundamental Plane (FP) of early-type galaxies in the clusters of galaxies \\clfort\\ at $z=1.013$. This is the first detailed FP investigation of cluster early-type galaxies at redshift $z=1$. The distant cluster galaxies follow a steeper FP relation compared to the local FP. The change in the slope of the FP can be interpreted as a mass-dependent evolution. To analyse in more detail the galaxy population in high redshift galaxy clusters at $0.8<z<1$, we combine our sample with a previous detailed spectroscopic study of 38 early-type galaxies in two distant galaxy clusters, RXJ0152.7-1357 at $z=0.83$ and RXJ1226.9+3332 at $z=0.89$. For all clusters Gemini/GMOS spectroscopy with high signal-to-noise and intermediate-resolution has been acquired to measure the internal kinematics and stellar populations of the galaxies. From HST/ACS imaging, surface brightness profiles, morphologies and structural parameters were derived for the galaxy sample. The least massive galaxies ($M=2\\times10^{10}M_{\\sun}$) in our sample have experienced their most recent major star formation burst at $z_{\\rm form}\\sim 1.1$. For massive galaxies ($M>2\\times10^{11}M_{\\sun}$) the bulk of their stellar populations have been formed earlier $z_{\\rm form}\\ga 1.6$. Our results confirm previous findings by J\\o rgensen et al. This suggests that the less massive galaxies in the distant clusters have much younger stellar populations than their more massive counterparts. One explanation is that low-mass cluster galaxies have experienced more extended star formation histories with more frequent bursts of star formation with shorter duration compared to the formation history of high-mass cluster galaxies.} ",
        "introduction": "Clusters of galaxies at intermediate-redshift ($1<z<2$) provide extremely useful laboratories to map the large-scale structure of the universe and study the formation and evolution of cluster galaxies.  There are only a small number of clusters of galaxies known at redshifts $z>1$ and our understanding of the galaxy population in these high redshift clusters is limited (e.g., Rosati et al. 2004; Demarco et al. 2007; Lidman et al. 2008; Mei et al. 2009). $X$-ray surveys (e.g., Stanford et al. 2001) have detected about a dozen clusters with $z\\ge 1$. Current surveys carried out in the $X$-rays (Romer et al.\\ 2001; Stanford et al. 2006; Finoguenov et al. 2007), optical (e.g., Gladders \\& Yee 2005), infrared (Stanford et al.\\ 2005; Lawrence et al.\\ 2007), and radio (Bran\\-chesi et al.\\ 2006) provide the basis for new but small number of discoveries. Upcoming surveys using the Sunyaev-Zel'dovich effect are expected to discover between 100 to 1000 new massive galaxy clusters at high-redshifts of $z>1$ (Carlstrom et al.\\ 2002).  In this context, new spectroscopic observations of clusters of galaxies play a key role for understanding galaxy evolution in dense environments. Previous studies at $z\\sim1$ concentrated on a few of the brightest (hence more massive $\\ga 2\\times10^{11}\\,M_{\\sun}$) early-type cluster members (van Dokkum \\&  Stanford 2003; Holden et al.\\ 2005), because high signal-to-noise ($S/N$) spectroscopy of distant galaxies is very expensive in telescope time ($\\ga$20 hours at 8-10m facilities). In particular, these two Fundamental Plane analyses were restricted to three and four brightest early-type cluster members, respectively.  At redshifts of $z\\sim1.5$, both cluster (e.g., Cucciati et al. 2006; Cooper et al. 2007) and field (e.g., Bell et al. 2004; Faber et al. 2007; Ferreras et al. 2009; Pozzetti et al. 2009) surveys show a significant increase (by about a factor of two) in the number density of early-type galaxies from high\\-er redshift to the present day. In other words, only about 50\\% of the total baryonic (stellar) mass in early-type galaxies has been generated before $z_{\\rm form}\\sim1.5$, whereas the remainder of the bulk of the stars has been assembled at much later epochs of the universe $z_{\\rm form}\\la 2$.  Through an investigation of the Fundamental Plane constraints on the evolution and formation history of early-type galaxies as well as their dark matter content are possible. The Fundamental Plane (FP) is a tight linear relation in three-dimensional log-space defined by the effective (half-light) radius $r_{{\\rm e}}$, the average surface brightness within $r_{{\\rm e}}$ ($\\langle I_{{\\rm e}}\\rangle$) and the central velocity dispersion ($\\sigma$) of early-type galaxies (Djorgovski \\& Davis 1987; Dressler et al.\\ 1987; J\\o rgensen et al.\\ 1996, hereafter JFK96). This scaling relation has proven to be a powerful tool in measuring the luminosity-weighted average ages of early-type galaxies, \\newline both at local (e.g., Bender et al.\\ 1992; JFK96; Bernardi et al.\\ 2003) and at intermediate up to redshifts of $z\\sim1.2$ (e.g., J\\o rgensen et al.\\ 1999; Wuyts et al.\\ 2004; Fritz et al.\\ 2005; di Serego Alighieri et al. 2005; Treu et al.\\ 2005; J\\o rgensen et al.\\ 2006, 2007; Fritz et al.\\ 2009). The mean age of the stellar populations in these galaxies can be directly probed, because an evolution in the zero-point offset of the FP with increasing redshift can be directly related to a change in the average mass-to-light ($M/L$) ratio of the galaxies under consideration.  At $z\\sim1$ our understanding of the FP is very limited. Previous works based on small number statistics at $z\\sim1.3$ found a FP relation with a large intrinsic scatter, twice as high as in the Coma cluster (van Dokkum \\&  Stanford 2003; Holden et al.\\ 2005), with the properties of these early-type cluster galaxies poorly understood. This larger intrinsic scatter is in disagreement with the FP for two galaxy clusters at $z=0.8-0.9$ (J\\o rgensen et al.\\ 2006, 2007). In this high redshift domain selection effects play an important role. Therefore, both large, well defined sample sizes as well as high quality data are needed to separate observational limitations (like sample selection or cosmic variance) from the underlying physical processes in these distant galaxies (see also discussion in Fritz et al. 2009).  We present the FP for 12 early-type galaxies at $z=1.013$ that are associated with the rich, massive and $X$-ray luminous cluster of galaxies \\clfort. The galaxy cluster has been observed using a combination of high $S/N$ GMOS spectroscopy at the 8m Gemini North telescope and deep imaging of the Advanced Camera for Surveys (ACS) onboard Hubble Space Telescope (HST). Our most distant sample reaches an absolute magnitude limit of $M_B\\sim -21$ mag (rest-frame), which makes this work less susceptible to selection effects. Our study of \\clfort\\  is part of the Gemini/HST Galaxy Cluster Project (J\\o rgensen et al.\\ 2005), an extensive observational program to investigate the evolution, star formation and chemical enrichment history of galaxies in rich galaxy clusters from $z=1$ to the present-day universe. Previous results on the Gemini/HST Galaxy Cluster Pro\\-ject have been published on the stellar populations of the cluster RXJ0152.7-1357 at $z=0.83$ (J\\o rgensen et al.\\ 2005), the stellar content and FP of RXJ0142.0+2131 at $z$$=$$0.28$ (Barr et al. 2005, 2006) and the FP of RXJ0152.7-1357 and RXJ1226.9+3332 at $z=0.89$ (J\\o rgensen et al.\\ 2006, 2007). Throughout this article we adopt a $\\Lambda$CDM cosmology for a flat low-density Universe with $H_0$$=$$70\\,{\\rm km\\,s^{-1}}$ ${\\rm Mpc^{-1}}$, $\\Omega_M=0.3$, and $\\Omega_{\\rm \\Lambda}=0.7$. Unless otherwise noted, all magnitudes are given in the Vega system. ",
        "conclusions": "Galaxies in clusters at intermediate redshift are useful pro\\-bes to study the formation and the evolution of early-type galaxies. By measuring the kinematics and structure parameters of these galaxies their evolution in mass and mass-to-light ($M/L$) ratios can be derived. This allows us to put constraints on the formation epoch and subsequent evolution of spheroidal galaxies up to the present-day as well as to critically test the models of galaxy formation and evolution.  In this paper we have presented the results of a detailed study of early-type galaxies in the distant clusters of galaxies \\clfort\\ at $z=1.013$. Based on HST/ACS imaging, surface brightness profiles, morphologies and stru\\-ctural parameters were derived for 30 galaxies within the HST/ACS field-of-view of the galaxy cluster. We have constructed the first detailed Fundamental Plane (FP) of 12 clu\\-ster early-type galaxies at redshift $z=1$. This analysis is based on a sample size that represents a factor of 3 or more improvement compared to previous studies. We have combined our distant cluster sample with our previous detailed spectroscopic study of 38 early-type galaxies in two distant galaxy clusters at $z=0.8-0.9$, RXJ0152.7-1357 and RXJ1226.9+3332 (J\\o rgensen et al.\\ 2006, 2007). This allows us to study in more detail the early-type galaxy population in high redshift galaxy clusters up to redshift unity. Compared to the local reference, the distant cluster galaxies follow a steeper FP relation and $M/L$-mass relation. This slope change can be interpreted as a difference in the epoch of the last major star formation episode for galaxies with different stellar masses. The least massive galaxies ($M=2\\times10^{10}M_{\\sun}$) in our sample have experienced their most recent major star formation burst at $z_{\\rm form}\\sim 1.1$, whereas for massive galaxies ($M>2\\times10^{11}M_{\\sun}$) the majority of their stellar populations have been formed earlier $z_{\\rm form}\\ga 1.6$. The dependence of the formation epoch on the mass confirms previous results by J\\o rgensen et al.\\ (2006, 2007).  A similar mass-dependent evolution with more extended SFH has recently been found for field early-type galaxies (Fritz et al. 2009). Our results can be understood if less massive cluster galaxies have on average younger stellar populations. The distant galaxy population in these low-mass galaxies could have been built up over longer time scales and over more extended star formation histories of shorter duration but more frequent star formation bursts, compared to the formation history of the high-mass galaxy cluster counterparts.    Our results support a galaxy formation scenario according to the \\textit{down-sizing} picture (Cowie et al. 1996), where the mass of galaxies hosting star formation processes decreases with the age of the Universe. A combined analysis of the FP and the absorption line indices for \\clfort\\ as well as RXJ0152.7-1357 and RXJ1226.9+3332 will be presented in J\\o rgensen et al.\\ (2009)."
    },
    "0910/0910.4013_arXiv.txt": {
        "abstract": "We use a high-resolution $N$-body simulation to investigate the influence of background galaxy properties, including redshift, size, shape and clustering, on the efficiency of forming giant arcs by gravitational lensing of rich galaxy clusters.  Two large sets of ray-tracing simulations are carried out for 10 massive clusters at two redshifts, i.e. $z_{\\rm l} \\sim 0.2$ and $0.3$. The virial mass ($M_{\\rm vir}$) of the simulated lens clusters at $z\\sim0.2$ ranges from $6.8\\times10^{14} h^{-1} {M_{\\odot}}$ to $1.1\\times 10^{15} h^{-1} M_{\\odot}$.  The information of background galaxies brighter than 25 magnitude in the $I$-band is taken from Cosmological Evolution Survey (COSMOS) imaging data. Around $1.7\\times 10^5$ strong lensing realizations with these images as background galaxies have been performed for each set. We find that the efficiency for forming giant arcs for $z_{\\rm l}=0.2$ clusters is broadly consistent with observations.  % Our study on control source samples shows that the number of giant arcs is decreased by a factor of 1.05 and 1.61 when the COSMOS redshift distribution of galaxies is adopted, compared to the cases where all the galaxies were assumed to be in a single source plane at $z_{\\rm l}=1.0$ and $z_{\\rm l}=1.5$, respectively. We find that the efficiency of producing giant arcs by rich clusters is weakly dependent on the source size and clustering. Our principal finding is that a small proportion ($\\sim 1/3$) of galaxies with elongated shapes (e.g. ellipticity $\\epsilon=1-b/a>0.5$) can boost the number of giant arcs substantially. Compared with recent studies where a uniform ellipticity distribution from 0 to 0.5 is used for the sources, the adoption of directly observed shape distribution increases the number of giant arcs by a factor of $\\sim2$. Our results indicate that it is necessary to account for source information and survey parameters (such as point-spread-function, seeing) to make correct predictions of giant arcs and further to constrain the cosmological parameters. ",
        "introduction": "Giant arcs are spectacular examples of strong gravitational lensing found in rich galaxy clusters. Background galaxies are stretched into long, thin arcs by the intense foreground gravitational field. Hundreds of giant arcs have been found in both optically-selected and X-ray selected clusters (e.g. \\citealt{Luppino99, ZG03, Gladders03, Sand05}).  Massive clusters are efficient producer of giant arcs and the number of giant arcs is a good indicator of the abundance of massive clusters. The halo mass function is very sensitive to the cosmological parameters, especially at the massive end. Moreover, the internal structures of massive clusters, such as substructures and ellipticity, also depend on the cosmological parameters. They all affect the lensing probability (optical depth, e.g., \\citealt{Mene03a, Mene03b, Wambsganss04, Torri04, Dalal04, Li05, OLS03, Oguri08, Hilbert07, Puchwein05, Puchwein09}) in various degree. Therefore, the observation of giant arcs is a useful probe of the cosmological model, in particular the matter power-spectrum normalization, $\\sigma_8$ \\citep{Li06b}, the matter density, $\\omega0$, and to a less extent the cosmological constant, $\\Omega_{\\Lambda,0}$.  However, in order to use the observations of giant arcs to constrain the cosmological models, one has to thoroughly understand how the lensing probability is affected by the distribution of background source galaxies, as well as the intrinsic properties of lens population. It has been shown that the lensing probability increases significantly with the increase of the source redshift, thus it is necessary to quantify the redshift distribution of source galaxies in lensing studies \\citep{Wambsganss04}. The lensing efficiency of massive clusters does not depend on the source size significantly \\citep{Li05}, although not all the real source information has been used, in particular their shapes.  \\citet{Horesh05} first adopted realistic galaxy images with known photometric redshifts from the Hubble Deep Field (HDF) as background sources. They selected clusters from a cosmological $N$-body simulation in a $\\sim140 h^ {-1}{\\rm Mpc}$ box at $z_{\\rm l}\\sim 0.2$. The mass of the simulated lens clusters ($0.54$--$1.1\\times10^{15} h^ {-1}{\\rm  M_{\\odot}}$) is similar to that of 10 X-ray--selected clusters \\citep{Smith05}, the (mean) mass range of which is around $6.3\\times10^{14} h^{-1} M_\\odot \\leq M_{200}\\leq2.0\\times10^{15} h^{-1} M_{\\odot}$. Note that the mass is calculated based on the X-ray luminosity by using the $L_{\\rm x}$-$M_{200}$ relation \\citep{Reiprich02}. They argued that the probability of producing giant arcs is $\\sim 1$ per cluster after observational effects are included, and emphasize that the number of giant arcs produced by the simulated clusters in their adopted cosmology (with $\\omega0=0.3, \\Omega_{\\Lambda,0}=0.7$ and the power-spectrum normalization of $\\sigma_8=0.9$) is consistent with the one observed in the massive clusters at redshift $0.171<z_c<0.255$.  In this paper, we focus on the impact of background galaxy properties, rather than the properties of the lens clusters, on the efficiency of producing giant arcs. We use the COSMOS data as direct input for the properties of background galaxies. We investigate how the strong lensing probabilities are affected by the shape, size, redshift distributions and clustering of background galaxies. Instead of adopting all these source properties as a whole \\citep{Horesh05}, we vary each of the source properties in turn to gain a clear understanding about their individual impacts on strong lensing statistics. The COSMOS data has a substantial advantage over the HDF as it covers a much larger sky area (its sky coverage is 1.6 square degrees vs. 5.3 square arcmins for the HDF) as used in \\citet{Horesh05}, and thus is less affected by the cosmic variance. While we confirm the effects of source redshift and size distributions, we find that the shape distribution of real galaxies has a dramatic effect on the lensing probability. With the COSMOS images, the lensing probability is boosted by a factor of $2$ relative to that based on simple assumptions of the shape distribution (e.g. random ellipticity between 0 and 0.5). Our results indicate that the shape of galaxies is an important factor in matching the theoretical predictions and the observations of giant arcs. We will also show that clustering of background galaxies has a negligible effect on the lensing efficiency of producing giant arcs.  This paper is structured as follows. In \\S2, we discuss COSMOS and background source population.  In \\S3, we discuss the simulation we use and our lensing methodology.  In \\S4, we present our main results and compare with recent studies. We finish in \\S5 with a short summary and a brief discussion of the implications of our results. ",
        "conclusions": "In this paper, aiming to study the impacts of background sources on strong lensing statistics, we use the $I$-band galaxy image data of HST/ACS in the COSMOS to quantify the distributions of background source size, shape and redshift. Each galaxy image is extracted by SExtractor limiting a surface brightness down to $\\sim 25 \\,{\\rm mag/arcsec^2}$. The redshift of each galaxy image is obtained by matching its celestial position with the COSMOS photometric catalog. The selected sources are then lensed by 10 massive clusters of the mass range $6.8 \\times 10^{14} h^{-1} M_{\\odot}\\leq\\rm M_{\\rm vir}\\leq1.1\\times 10^{15} h^{-1} M_{\\odot}$ at $z_{\\rm l} \\sim 0.2$ and $6.0 \\times 10^{14} h^{-1} M_{\\odot}\\leq\\rm M_{\\rm vir}\\leq 8.1\\times 10^{14} h^{-1} M_{\\odot}$ at $z_{\\rm l} \\sim 0.3$ as the lensing clusters chosen from a cosmological simulation in the $\\Lambda$CDM cosmology ($\\Omega_{\\rm m,0}=0.268, \\Omega_{\\Lambda,0}=0.732$). 575 source tiles within the COSMOS field of around 2 ${\\rm deg^2}$ are fully used for the statistic study. 172500 ray-tracing lensing simulations are carried out for 10 lens clusters to reduce the statistical fluctuation at the expected lens redshifts (i.e. $z_{\\rm l}=0.2$ and 0.3). The incidence of giant-arc production in our simulation is roughly consistent with that observed in \\citet{Smith05}, after the density difference is taken into account. The impacts of source size, shape and redshift on strong lensing statistics are investigated in detail.  We find that the source size (less than a factor of 1.5) and clustering only have small effects on the production of giant arcs. In contrast, the dependencies on the source redshift and ellipticity are much more significant. The first was highlighted by \\citet{Wambsganss04}, while the second is the main new finding of the current work. We find that adopting the empirical ellipticity distribution of COSMOS increases the lensing probability by a factor of 2 (see Fig.~\\ref{fig:ze} and also the number ratio between Case 1 and 6). The boosting effect of the ellipticity of the background galaxies has not been emphasized in the previous works. It should be included in theoretical modeling of giant arcs in future. This may also be of particular relevance for setting constraints on the power-spectrum normalization parameter $\\sigma_8$.  There are a number of limitations in the present work. When scaling our results with the source number density, we implicitly assume that other source properties (like ellipticity and redshift) are the same in COSMOS as in a deeper observation (like HDF). However, ellipticity distribution of galaxies in a deeper survey may turn out to be more elliptical (\\citealt{Vincent05}; see also the discussion on the surface brightness limit), which would increase the boosting factor of source ellipticity on strong lensing efficiency. We would also have a larger fraction of galaxies at the high redshift tail end in a galaxy survey like HDF, which increases the mean strong lensing cross section. Therefore, both effects will  produce more giant arcs. For direct comparison with observation, the observational effects, such as specific instrumental point-spread-function and observational seeing, are not fully included in our lensing simulations. As mentioned in \\S\\ref{sec:seeing}, if the seeing was $1''$, then there may be only around $40\\%$ of giant arcs observed by a typical ground-based telescope compared with that by space-based one, such as HST.  The same shape measurement method is used for quantifying the ellipticity of the original galaxies and lensed images. It could bring in some inaccuracy for measuring the ellipticity of very round images, since it is mainly designed for quantifying arcs. Nevertheless, this effect is rather small and negligible to our results, since we mainly focus on the elliptical galaxies. Besides, the intrinsic pixelization will affect the ellipticity quantification of small sources, especially for sources of $D_{\\rm eff}<0.3''$ (see Fig.~\\ref{fig:slw}). The influence of pixelization would probably shift down the measured ellipticity because of its relative larger effect on width measurement.  Since the number of such tiny sources is small ($\\sim 5\\%$) and the pixelization does not significantly change the ellipticity value on average (much smaller than the effect of seeing), this influence would not be important to our analysis.  According to \\cite{Torri04}, cluster substructures are important for strong lensing efficiency in a major merger case, especially at the early state. As we adopt a softening length of 30 $h^{-1}\\rm{kpc}$, the softening mainly takes effect in the center of major merging clumps, while it would also smooth away the small substructures of a similar scale in a minor merger case or those remaining in a main halo. The change of the lensing cross section due to this kind of smoothing would not be significant in both cases. Besides, according to a comparison plot of optical depth in Fig.~4 of \\cite{Li05}, the result agrees within $25\\%$ of \\cite{Wambsganss04}, in which they adopt a much smaller softening length of $3.2 h^{-1}\\rm{kpc}$ than 30 $h^{-1}\\rm{kpc}$ in our case, for a source redshift of $z_{\\rm s}=1.0$. Therefore, although the relatively large value of softening length could reduce the lensing efficiency, it would not affect our results significantly. Moreover, lensing clusters are selected from a dark matter only simulation, therefore, we cannot investigate the strong lensing dependence on source properties for clusters with realistic baryon distributions.  We have only used one surface brightness limit (e.g., a detection threshold of 1.5 $\\sigma$) to extract the background galaxy images. Since the redshift of galaxy is mainly constrained by the matching with photometric catalog, it will not change the distribution by using a higher detection threshold in our case. A variation of detection threshold would change the galaxy size, but the dependence of lensing efficiency on source size is weak (see \\S\\ref{sec:sourcesize}), thus the main results should remain unchanged. The source shape distribution also depends on the surface brightness threshold. Indeed, at a lower surface brightness limit, the shape of the galaxies tends to more elliptical \\citep{Vincent05}. Thus at fainter surface brightness, the shape influence is expected to play a more important role on giant-arc production for the same lens cluster population. It is meaningful to do ray-tracing experiments by using several surface brightness limits to quantify the shape impact on strong lensing in different survey depths.  The different mass range of the simulated clusters and X-ray selected ones could bring in quantitative uncertainties in the giant-arc statistics, since the strong lensing cross section strongly depend on mass, e.g., the mean cross section changes by nearly 4 orders of magnitude over mass ranging from $10^{14}$ $h^{-1}{M_{\\odot}}$ to $\\sim10^{15}$ $h^{-1}{M_{\\odot}}$ in Fig.~7 of \\cite{Hen07a}. However, the mean cross section ratio is less than 1.5 between the observed and our simulated clusters, according to the plot. Moreover, they use a larger mass density parameter $\\omega0=0.3$ and a higher matter power-spectrum normalization $\\sigma_8=0.95$ in their simulation. Besides, the simulated clusters are put at a much  higher redshift $z_{\\rm l}=0.41$. All these effects help to make the mass dependence of lensing probability stronger. Therefore, the difference of mean cross section in two comparing samples would be much less than a factor of 1.5 and becomes negligible in this work. Nevertheless, it is still necessary and interesting to select more comparable (between the simulated and X-ray selected) lens cluster samples in mass and redshift range for a fair comparison of giant-arc producing efficiency in the future."
    },
    "0910/0910.3633_arXiv.txt": {
        "abstract": "In strong magnetic fields the transport coefficients of strange quark matter become anisotropic. We determine the general form of the complete set of transport coefficients in the presence of a strong magnetic field. By using a local linear response method, we calculate explicitly the bulk viscosities $\\zperp$ and $\\zpara$ transverse and parallel to the $B$-field respectively, which arise due to the non-leptonic weak processes $u+s\\leftrightarrow u+d$. We find that for magnetic fields $B<10^{17}$ G, the dependence of $\\zperp$ and $\\zpara$ on the field is weak, and they can be approximated by the bulk viscosity for zero magnetic field. For fields $B>10^{18}$ G, the dependence of both $\\zperp$ and $\\zpara$ on the field is strong, and they exhibit de Haas-van Alphen-type oscillations. With increasing magnetic field, the amplitude of these oscillations increases, which eventually leads to negative $\\zperp$ in some regions of parameter space. We show that the change of sign of $\\zperp$ signals a hydrodynamic instability. As an application, we discuss the effects of the new bulk viscosities on the r-mode instability in rotating strange quark stars.  We find that the instability region in strange quark stars is affected when the magnetic fields exceeds the value $B= 10^{17}$ G. For fields which are larger by an order of magnitude, the instability region is significantly enlarged, making magnetized strange stars more susceptible to $r$-mode instability than their unmagnetized counterparts. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.1894_arXiv.txt": {
        "abstract": "We present results of optical spectroscopic and $BVR_\\mathrm{C}I_\\mathrm{C}$ photometric observations of 77 pre-main sequence (PMS) stars in the Cepheus flare region. A total of 64 of these are newly confirmed PMS stars, originally selected from various published candidate lists. We estimate effective temperatures and luminosities for the PMS stars, and comparing the results with pre-main sequence evolutionary models we estimate stellar masses of 0.2--2.4\\,M$_{\\sun}$ and stellar ages of 0.1--15~Myr. Among the PMS stars, we identify 15 visual binaries with separations of 2--10~arcsec. From archival IRAS, 2MASS, and Spitzer data, we construct their spectral energy distributions and classify 5\\% of the stars as Class~I, 10\\% as Flat SED, 60\\% as Class~II, and 3\\% as Class~III young stellar objects (YSOs). We identify 12 CTTS and 2 WTTS as members of NGC~7023, with mean age of 1.6~Myr. The 13 PMS stars associated with L\\,1228 belong to three small aggregates: RNO~129, L\\,1228\\,A, and L\\,1228\\,S. The age distribution of the 17 PMS stars associated with L\\,1251 suggests that star formation has propagated with the expansion of the Cepheus flare shell. We detect sparse aggregates of $\\sim$6--7~Myr-old PMS stars around the dark clouds L\\,1177 and L\\,1219, at a distance of $\\sim 400$\\,pc. Three T~Tauri stars appear to be associated with the Herbig Ae star SV~Cep at a distance of 600~pc. Our results confirm that the molecular complex in the Cepheus flare region contains clouds of various distances and star forming histories. ",
        "introduction": "\\label{Sect_1}   Star forming regions in our Galactic environment offer a unique opportunity for understanding the birth of planetary systems. There is growing evidence that not only the stellar content and structure of young stellar groups, but also the evolution of protostellar envelopes and accretion disks depend on the star forming environment \\citep[e.g.][]{Hester05,Dullemond06}. Cross-comparing various star forming environments is important for understanding the origins and early evolution of stars and planets.  In this paper we present our results of spectroscopic and photometric studies of a population of pre-main sequence stars associated with the molecular clouds of the Cepheus flare, located at 200--400~pc from the Sun.  A detailed review on the ISM and young stellar populations in the Cepheus flare region can be found in \\citet*{Kun08}. Here we briefly review the highlights that are relevant for this study. The Cepheus flare \\citep{Hubble} molecular cloud complex consists of a large amount of dense ISM between $100\\degr \\leq l \\leq 120\\degr$ and above $b \\geq 10\\degr$ \\citep*{LyndsD,TDS, CB88,DUK}.  Though the CO maps of the region presented by \\citet{Lebrun} and \\citet{GLADT} indicate a coherent cloud complex, there is evidence for multiple cloud layers over the Cepheus flare region.  For instance, stars illuminating reflection nebulae or displaying 100\\,\\micron \\ excess due to their dusty environment, can be found at various distances \\citep*[e.g.][]{Racine,KVSz}, suggesting the presence of more than one dust layer between about 200 and 800~pc. Wolf diagrams of the region indicate extinction layers at 200, 300, and about 450~pc \\citep{Kun98}.  Radial velocity measurements of the interstellar matter suggest the presence of several gas layers \\citep[e.g.][]{Heiles,YDMOF}.  \\citet{Olano} modeled the space distribution and kinematics of the interstellar neutral and molecular gas by an expanding shell, centered on $(l,b)\\approx (120\\degr,+17\\degr)$, at a distance of about 300~pc. The presence of the {\\em Cepheus flare Shell\\/} (CFS), an old supernova remnant, readily explains the divergent cloud distances and the wide range of radial velocities. The giant radio continuum feature {\\em Loop~III} \\citep{Berk73, Spoelstra}, conspicuous in the {\\it WMAP\\/} K-band polarization map \\citep{Page07}, is nearly concentric with the molecular shell. Several giant far-infrared loops, possibly old supernova shells, have been identified in the IRAS images of the Cepheus flare region \\citep*{KMT}.  Signposts of low-mass star formation in the Cepheus flare region have been detected in the molecular clouds L\\,1082 (B\\,150, GF\\,9), L\\,1148/L\\,1158, L\\,1172/L1174 (associated with the young cluster NGC\\,7023), L\\,1228, L\\,1177 (CB 230), L\\,1219 (B\\,175), L\\,1221, L\\,1251, and L\\,1261 (CB~244) \\citep[see review of ][]{Kun08}. Recently, \\citet{Kirk09} presented a list of 133 young stellar objects and candidates, based on Spitzer observations of L\\,1148/L\\,1158, L\\,1172/L1174, L\\,1228, L\\,1221, and L\\,1251. Independent distance estimations of individual clouds suggest various distances between 150 and 500~pc, but are clustered around 300\\,pc. The young open cluster NGC\\,7129, associated with the dark cloud L\\,1181, is projected on the Cepheus flare region at ($l$,$b$)=(105.4,+9.9), but is more distant than the bulk of clouds (800--1250~pc, see \\citet{Kun08}. The central region of the complex is lacking signatures of active star formation. The physical connection between the star-forming clouds remains rather uncertain.  Scarce information is available on the pre-main-sequence (PMS) stellar population of the region. \\citet{HBC} list only 14 young objects in the Cepheus flare region: RNO\\,124 and PV Cep (associated with the cloud complex L\\,1148/L\\,1158), HD\\,200775, EH Cep, FU Cep, FV Cep (members of NGC\\,7023), Lk\\ha~234, BD\\,+65\\degr1637, V350~Cep, [SVS76]~NGC\\,7129~6 (members of NGC\\,7129), AS~507 (associated with L\\,1261), and the apparently isolated Herbig Ae/Be stars BH Cep, BO Cep, and SV Cep. Six further stars in NGC\\,7129 (MEG\\,1--MEG\\,3, [HL85]\\,14: \\citealt{MEG93}; GGD~33a: \\citealt{MEFG94}, and V391~Cep: \\citealt{Semkov93}) show T Tauri-like emission spectra, but no spectral types have been given for them in the literature. \\citet{Tachihara05} identified 16 weak-line T~Tauri stars in the region, by spectroscopic and photometric follow-up studies of X-ray emitting stars found in the ROSAT database. These stars are not obviously associated with any dark cloud. Understanding star formation in each dark cloud and relating star formation rates across the dark clouds requires a comprehensive deep survey.  Here, we use optical and infrared observations to develop a better picture of star formation in this region. To constrain the nature of the candidate young stars, we combine new optical spectroscopy and BVRI photometry with archival infrared data from IRAS, 2MASS, and Spitzer.  We describe our observations and data reduction in Sect.~\\ref{Sect_obs}. We use these data to estimate ages, luminosities, masses, and disk properties of individual young stars and to analyze the distributions of these physical parameters among each cloud. Section~\\ref{Sect_res} contains the main results, namely the HR diagrams and spectral energy distributions of the target stars. We discuss the implications of our results for the star formation history of the region in Sect.~\\ref{Sect_disc}. ",
        "conclusions": "From spectroscopic and photometric observations, we estimated the effective temperatures and luminosities of 78 PMS stars in the Cepheus flare region, 65 of which are newly confirmed PMS stars. Comparing the results with pre-main sequence evolutionary models we estimate stellar masses of 0.2--2.4\\,M$_{\\sun}$ and stellar ages of 0.1--15~Myr. From archival IRAS, 2MASS, and Spitzer data, we constructed the spectral energy distributions; 5\\% of the stars are Class~I, 5\\% Flat, 81\\% Class~II, and 9\\% Class~III.  Star forming clouds and PMS stars of the region belong to the complexes as follows. \\begin{itemize} \\item Star formation associated with the large expanding system of Cepheus flare Shell/Loop~III occurs in the clouds L\\,1148/L\\,1158, L\\,1172/L\\,1174, L\\,1251, L\\,1228 and L\\,1261/L\\,1262. The latter two clouds are located on the approaching side of the expanding shell. \\item A sparse association consisting of 6--7~Myr old PMS stars at the low-latitude part of the region ($105\\degr \\lesssim l \\lesssim 112\\degr$, $10\\degr\\lesssim b \\lesssim 14\\degr$) can be found at a distance of 380--400\\,pc. The  shapes of the clouds in this part of the region suggest shocks from the direction of the Galactic plane. \\item Low-mass PMS stars are associated with the Herbig~Ae star SV~Cep, located at $\\sim$600\\,pc. \\item NGC\\,7129, at $\\sim$800\\,pc from the Sun, is probably associated with the Cepheus Bubble, created by the older subgroup of the association Cep~OB2. Stars have been formed in at least two clumps of its natal cloud L\\,1181. \\end{itemize} The A1-type \\ha\\ emission star associated with IRAS~22219+7901 is probably more distant than the Cepheus flare clouds. Its nature remains uncertain."
    },
    "0910/0910.3069_arXiv.txt": {
        "abstract": "Two recent papers (Ghez et al. 2008, Gillessen et al. 2009) have estimated the mass of and the distance to the massive black hole in the center of the Milky Way using stellar orbits. The two astrometric data sets are independent and yielded consistent results, even though the measured positions do not match when simply overplotting the two sets. In this letter we show that the two sets can be brought to excellent agreement with each other when allowing for a small offset in the definition of the reference frame of the two data sets. The required offsets in the coordinates and velocities of the origin of the reference frames are consistent with the uncertainties given in Ghez et al. (2008). The so combined data set allows for a moderate improvement of the statistical errors of mass of and distance to Sgr~A*, but the overall accuracies of these numbers are dominated by systematic errors and the long-term calibration of the reference frame. We obtain $R_0 = 8.28 \\pm 0.15 |_\\mathrm{stat} \\pm 0.29|_\\mathrm{sys}\\,$kpc and $M_\\mathrm{MBH} = 4.30 \\pm 0.20 |_\\mathrm{stat} \\pm 0.30|_\\mathrm{sys} \\times 10^6\\,M_\\odot$ as best estimates from a multi-star fit. ",
        "introduction": "The motions of stars in the immediate vicinity of Sgr~A* have been tracked since 1992 at the NTT/VLT and since 1995 at the Keck telescope \\citep{Eckart:1996p163,Ghez:1998p118}. With the detection of accelerations \\citep{Ghez:2000p138} and the determination of the first orbit \\citep{Schodel:2002p153} these measurements provided firm evidence for the existence of a massive black hole (MBH) at the center of the Milky Way, coincident with the radio-source Sgr~A*. The stars are used as test particles for the gravitational potential in which they move. Particularly important is the bright ($m_\\mathrm{K}\\approx14$) star S2 (S0-2 in the Keck nomenclature) orbiting the MBH in 15.9 years on an ellipse with an apparent diameter of $\\,\\approx0.2''$. Together with radial velocity data, the astrometric data allow for a geometric determination of $R_0$, the distance to the Galactic Center \\citep{Salim:1999p1362,Eisenhauer:2003p128}. The number of orbits has increased to more than 20 since then \\citep{Ghez:2005p1681,Eisenhauer:2005p117,Gillessen:2009p1117} and in particular the full S2 orbit has recently been used for improved determinations of the mass and of $R_0$ \\citep{Ghez:2008p945,Gillessen:2009p1117}.  These two astrometric data sets are independent since they were obtained at different telescopes and were analyzed by different teams using different tools. On the other hand, the radial velocity data attached to the astrometric data sets largely overlap, mainly because this technically is straight-forward. The radial velocities refer to the local standard of rest and hence the inclusion of other data into a given data set needs no special care. The situation is different for astrometric data, for which only an approximate realization of an absolute reference frame currently is possible \\citep{Reid:2007p169}. This means that the exact definition of coordinates is a matter of the respective data analysis, and simply merging two lists of astrometric positions will fail (see figure~\\ref{f1}).  The two groups \\citep{Ghez:2008p945,Gillessen:2009p1117} come to very similar conclusions for the mass of Sgr~A* and $R_0$. From VLT data \\cite{Gillessen:2009p1117} derived (when using the S2 orbit only) \\begin{eqnarray} R_0&=&8.31 \\pm 0.33|_\\mathrm{stat} \\,\\mathrm{kpc}\\nonumber\\\\ M_\\mathrm{MBH} &=& 4.29 \\pm 0.07|_\\mathrm{stat} \\pm 0.34|_{R_0}\\,M_\\odot\\,\\, \\end{eqnarray} where the statistical error on the mass is for a fixed distance and the second error term is due to the statistical error on $R_0$. The latter includes already the coordinate system uncertainty. In addition, a systematic error of $\\Delta R_0 \\approx \\pm^{\\,0.5}_{\\,1.0}\\,$kpc is present, owing to the uncertainties of a) how much the 2002 data (close to periastron) can be trusted and b) the assumption that the effective potential is Keplerian. The systematic error of $R_0$ also influences $M_\\mathrm{MBH}$ since the two quantities are correlated. \\cite{Ghez:2008p945} obtained from the Keck data set (neglecting 2002 data, and fixing the radial velocity $v_z$ of the MBH to 0) \\begin{eqnarray} R_0&=&8.4 \\pm 0.4\\,\\mathrm{kpc} \\nonumber\\\\ M_\\mathrm{MBH} &=& 4.5 \\pm 0.4\\,M_\\odot\\,\\,. \\end{eqnarray} The errors are of statistical nature and the mass error includes the uncertainty induced by the error on $R_0$. Probably, also a systematic error should be taken into account here. Leaving $v_z$ free yielded somewhat larger statistical errors, e.g. $\\Delta R_0=0.6\\,\\mathrm{kpc}$.  The aim of this letter is to investigate the combined data set. We show that not only the fit results are in excellent agreement, but also that the combined data can be described very well by a single Keplerian orbit. The such combined data set also yields an improvement in the statistical fit errors for the derived parameters. This can be understood qualitatively by inspecting the data: Before 2002, the Keck measurements appear superior to our NTT/VLT data set. At that time, our group was using the NTT with an aperture of $3.6\\,$m, which is notably smaller than $10\\,$m for the Keck telescopes. After 2002, the VLT data has comparable errors to the Keck data for individual data points but a much denser time sampling. This is owed to the location on the southern hemisphere of the VLT and the generous allocation of telescope time to the Galactic Center projects. Hence, by suitably combining the two data sets, one should obtain a data set which combines the respective advantages.\\footnote{We use data collected between 1992 and 2009 at the European Southern Observatory, both on Paranal and La Silla, Chile; in particular from the Large Programs 179.B-0261 and 183.B-0100.} ",
        "conclusions": "Since more than 10 years two groups have assembled independently astrometric data sets for the stars in the central arcsecond of the Milky Way. The data sets are truly independent. They were obtained at different telescopes with different instruments. The analysis chains used different tools (deconvolution in the NTT/VLT case, PSF-fitting for the Keck data), and the definition of astrometric coordinates is implemented in different ways. While it was reassuring, that \\cite{Ghez:2008p945} and \\cite{Gillessen:2009p1117} concluded very similarly for $M_\\mathrm{MBH}$ and $R_0$, it was not clear that the agreement actually also holds for the underlying data. We were able to show that indeed the most simple assumption - namely that the two data sets only differ in the coordinate system definition - is sufficient to perfectly map the two data sets on top of each other.  We presented for the first time a combined orbit fit, which however only moderately improves the accuracy by which $M_\\mathrm{MBH}$ and $R_0$ can be derived from the S2 data: \\begin{eqnarray} R_0 &=& 8.34 \\pm 0.27 |_\\mathrm{stat} \\pm 0.5|_\\mathrm{sys}\\,kpc \\nonumber\\\\ M_\\mathrm{MBH} &=& 4.40 \\pm 0.27 |_\\mathrm{stat} \\pm 0.5|_\\mathrm{sys} \\times 10^6\\,M_\\odot\\,\\,. \\end{eqnarray} For the sake of completeness, we cite also the updated numbers for a fit using not just the combined S2 data, but in addition S1, S8, S12, S13 and S14, the same stars \\cite{Gillessen:2009p1117} had used: \\begin{eqnarray} R_0 &=& 8.28 \\pm 0.15 |_\\mathrm{stat} \\pm 0.29|_\\mathrm{sys}\\,kpc \\nonumber\\\\ M_\\mathrm{MBH} &=& 4.30 \\pm 0.20 |_\\mathrm{stat} \\pm 0.30|_\\mathrm{sys} \\times 10^6\\,M_\\odot\\,\\,. \\end{eqnarray}\\\\ Besides the expected improvement in $\\Delta R_0|_\\mathrm{stat}$ from $0.17\\,$kpc to $0.15\\,$kpc, also the systematic errors have decreased mildly compared to the previous work due to the smaller influence of the S2 2002 data.  We conclude that the combination does not help much in overcoming current limitations. The true value of the two data sets actually is their independence, allowing for cross checks and lending credibility to the results.  Further substantial improvements in measuring the gravitational potential from Sgr~A* by means of stellar orbits probably will come from improved astrometry on existing data, from longer time lines with the existing instruments and finally from advances in instrumentation."
    },
    "0910/0910.4964_arXiv.txt": {
        "abstract": "% The relation between the globular cluster luminosity function (GCLF, ${\\rm d}N/{\\rm d}\\log{L}$) and globular cluster mass function (GCMF, ${\\rm d}N/{\\rm d}\\log{M}$) is considered. Due to low-mass star depletion, dissolving GCs have mass-to-light ($M/L$) ratios that are lower than expected from their metallicities. This has been shown to lead to an $M/L$ ratio that increases with GC mass and luminosity. We model the GCLF and GCMF and show that the power law slopes inherently differ (1.0 versus 0.7, respectively) when accounting for the variability of $M/L$. The observed GCLF is found to be consistent with a Schechter-type initial cluster mass function and a mass-dependent mass-loss rate. ",
        "introduction": "Even though globular cluster systems (GCSs) are considered to have formed during mergers, the shape of the globular cluster luminosity function (GCLF, ${\\rm d}N/{\\rm d}\\log{L}$) differs fundamentally from that of young massive clusters (YMCs) in merging galaxies. The luminosity function of YMCs follows a power law with index $-2$ over the full mass range down to a few $100~\\msun$, while the GCLF peaks at $\\sim 10^5~\\lsun$. This has been attributed to the ongoing tidal disruption of globular clusters (GCs) and the resulting destruction of low-mass, faint GCs \\citep{fall01,vesperini01}. In studies considering this mechanism, the mass-loss rate of GCs is considered to be independent of their mass (corresponding to a disruption time $t_{\\rm dis}\\propto M$) and the mass-to-light ratio ($M/L$) is assumed to be constant \\citep[e.g.,][]{fall01,jordan07}.  In order to trace back the merger history of galaxies using their globular cluster mass function (GCMF, ${\\rm d}N/{\\rm d}\\log{M}$), it is essential to obtain an accurate description for its evolution. Although the above studies reproduced the peaked shape of the GCMF, they did not account for two aspects of GC evolution: \\begin{itemize} \\item[(1)] The mass-loss rate of GCs does depend on cluster mass \\citep{baumgardt01,baumgardt03,lamers05,larsen09}, corresponding to $t_{\\rm dis}\\propto M^\\gamma$ with $\\gamma\\sim 0.7$ (see Eq.~\\ref{eq:dmdt}). This is due to the nonlinear scaling of the disruption time with the half-mass relaxation time \\citep[$t_{\\rm dis}\\propto t_{\\rm rh}^{0.75}$,][]{portegieszwart98,fukushige00}. A lower value of $\\gamma$ means that the dissolution rate of low-mass clusters is slowed down with respect to massive clusters. \\item[(2)] The $M/L$ ratio of GCs is not constant \\citep{mandushev91,baumgardt03,kruijssen08,kruijssen09} due to the preferential ejection of faint, low-mass stars from dissolving GCs. Because low-mass GCs have on average lost a larger fraction of their initial masses, $M/L$ increases with mass or luminosity. \\end{itemize} The effect of these aspects of GC evolution on the relation between the GCLF and the GCMF has been considered by \\citet{kruijssen09b}. The slope of the disruption-dominated low-mass side of the GCMF is always equal to $\\gamma$ \\citep{fall01,lamers05}. A mass-dependent mass-loss rate ($\\gamma=0.7$) therefore yields a different GCMF slope than a mass-independent mass-loss rate ($\\gamma=1$). The observed slope of the GCLF $\\sim 1$, which is seemingly consistent with the latter if a constant $M/L$ ratio is assumed (see Fig.~\\ref{fig:histcan}). However, the trend of increasing $M/L$ ratio with luminosity implies that this is not necessarily true because the slopes of the GCLF and the GCMF fundamentally differ. \\citet{kruijssen09b} have shown that the observed GCLF is in fact consistent with a GCMF with a low-mass slope of $\\sim 0.7$ when accounting for the trend of $M/L$ ratio with luminosity. \\begin{figure} \\center\\resizebox{12.4cm}{!}{\\plottwo{kruijssenj3fig1a.ps}{kruijssenj3fig1b.ps}} \\caption[]{\\label{fig:histcan} {\\it Left}: The observed GCLF of Galactic GCs \\citep{harris96}. {\\it Right}: The inferred GCMF of Galactic GCs (histogram) using $M/L_V=3$ \\citep[as in][]{fall01}. Overplotted is our model MF with a mass-dependent mass-loss rate (solid line, see Section~\\ref{sec:model}) adopting a dissolution timescale $t_0=1.6$~Myr. The dashed line shows the model for a cluster mass-independent mass-loss rate \\citep[as in][]{fall01}. Error bars are $1\\sigma$ Poissonian.} \\end{figure}  We revisit the calculations of \\citet{kruijssen09b} and include a more detailed model for the evolution of the stellar mass function (SMF) within a star cluster. Previously, the low-mass star depletion was approximated by increasing the lower stellar mass limit of the SMF. In reality, the SMF evolves more gradually. This has been included in new models of the evolution of the SMF in dissolving star clusters \\citep{kruijssen09c}, which are used here. ",
        "conclusions": "Using a realistic model for the evolution of the SMF in GCs, we have shown that the slopes of the GCLF and the GCMF differ. As GCs dissolve, their masses decrease more rapidly than their luminosities due to the ejection of faint, low-mass stars. As a result, the slope of the GCLF is $\\sim 1$ for luminosities below the peak luminosity, while the slope of the GCMF is $\\sim 0.7$ for masses below the peak mass. This is consistent with a mass-loss rate that depends on cluster mass and on the environment.  The new cluster models \\citep{kruijssen09c}, in which we account for the changing slope of the stellar mass function rather than shifting the lower stellar mass limit, show that the results from \\citet{kruijssen09b} also hold when more detailed models are applied. This substantiates that care should be taken when comparing GCLFs and GCMFs. Because of the relatively low number of GCs for which dynamical masses have been determined, we argue that the best method to compare them would be to model the GCLF rather than the GCMF, while accounting for the variability of the $M/L$ ratio. This would allow for a more accurate interpretation of the GCLF when studying correlations between its properties and the history of its galaxy."
    },
    "0910/0910.3858_arXiv.txt": {
        "abstract": "{Flow-induced instabilities of magnetic flux tubes are relevant to the storage of magnetic flux in the interiors of stars with outer convection zones. The stability of magnetic fields in stellar interiors is of importance to the generation and transport of solar and stellar magnetic fields. } {We consider the effects of material flows on the dynamics of toroidal magnetic flux tubes located close to the base of the solar convection zone, initially within the overshoot region. The problem is to find the physical conditions in which magnetic flux can be stored for periods comparable to the dynamo amplification time, which is of the order of a few years.} {We carry out nonlinear numerical simulations to investigate the stability and dynamics of thin flux tubes subject to perpendicular and longitudinal flows. We compare the simulations with the results of simplified analytical approximations.} {The longitudinal flow instability induced by the aerodynamic drag force is nonlinear in the sense that the growth rate depends on the perturbation amplitude. This result is consistent with the predictions of linear theory. Numerical simulations without friction show that nonlinear Parker instability can be triggered below the linear threshold of the field strength, when the difference in superadiabaticity along the tube is sufficiently large. A localised downflow acting on a toroidal tube in the overshoot region leads to instability depending on the parameters describing the flow, as well as the magnetic field strength. We determined ranges of the flow parameters for which a linearly Parker-stable magnetic flux tube is stored in the middle of the overshoot region for a period comparable to the dynamo amplification time. } {The longitudinal flow instability driven by frictional interaction of a flux tube with its surroundings is relevant to determining the storage time of magnetic flux in the solar overshoot region. The residence time for magnetic flux tubes with $2\\times 10^{21}$Mx in the convective overshoot layer is comparable to the dynamo amplification time, provided that the average speed and the duration of the downflow do not exceed about 50~m~s~$^{-1}$ and 100~days, respectively, and that the lateral extension of the flow is smaller than about $10^\\circ$. } ",
        "introduction": "\\label{sec:intro} Observations of large solar active regions are indicative of an organised subsurface magnetic field along the azimuthal (east-west) direction, with opposite polarity orientations (from east to west or vice versa) in the northern and southern hemispheres. Emerging magnetic structures are in a filamentary state, in the form of magnetic flux concentrations (flux tubes) of various sizes (e.g., sunspots, pores). Flux tubes rise in the convection zone, emerge at the surface, and form bipolar magnetic regions, which follow polarity rules (Hale's law) and systematic tilt angles (Joy's law).  Theoretical studies indicate that weak magnetic fields are transported by flux expulsion \\citep{msch84} and convective pumping \\citep{tobias01} to the lower boundary of the convection zone, where the toroidal magnetic field is amplified by radial and latitudinal velocity shear in the solar tachocline \\citep[see e.g.,][]{sis06}. The stably stratified lower convection zone, in particular the convective overshoot layer, is a likely location for the generation and storage of the large-scale azimuthal magnetic flux. Numerical studies simulating a layer of horizontal magnetic field in the bottom of the solar convection zone indicate that an initially uniform field underlying a field-free layer leads to the formation of magnetic flux tubes by magnetic Rayleigh-Taylor instability \\citep[e.g.,][]{fan01}. Following its formation, a toroidal magnetic flux tube can reach a mechanical equilibrium close to the bottom of the convection zone \\citep{mi92}. The emergence of magnetic flux tubes driven by magnetic buoyancy and the properties of active regions (low-latitude emergence, tilt angles, proper motions of sunspots) require azimuthal flux densities of the order of $10^5$~G in the overshoot region \\citep[][]{dsilva93,fan94,cale95,msch96}. The possibilities for the generation of these flux densities have been reviewed by \\citet{schrempel02}, \\citet{mschafm03}, and \\citet{afm07}.  The average duration of the solar activity cycle is about 11 years and the total magnetic flux emerging within one cycle ranges between orders of $10^{24}-10^{25}$~Mx. The amplification of the toroidal magnetic field to its maximum strength in about half an activity cycle raises the following question: how can the toroidal flux be stored stably for at least a few years in the dynamo amplification region? A related problem concerns the feedback of magnetic flux loss on the rate of toroidal flux generation. To more clearly understand magnetic flux generation and storage in the convective overshoot region, it is important to determine and constrain the effects of flows on the stability and dynamics of magnetic flux tubes. Here, we focus on: (a) the nonlinear development of flow-induced flux tube instabilities, which can be dynamically significant in the course of toroidal field amplification; and (b) effects of perpendicular flows on the storage of toroidal flux tubes. Consequently, some of the flow properties prevailing in the overshoot region can also be constrained, by requiring that flux tubes with field strengths of up to a few times $10^4$~G are stored in the overshoot region for about a few years.  We consider the possibility that the friction-induced instability \\citep[][]{pap2,pap3} leads to flux loss from the overshoot layer for field strengths below the Parker instability limit. We refer to such Parker-stable flux tubes as ``sub-critical'' throughout the paper. We investigate effects of flows on magnetic flux tubes in mechanical equilibrium to obtain quantitative estimates, and to answer the following question: how strongly do ($i$) finite-amplitude perturbations of flux tubes by perpendicular flows \\citep[see also][paper~I]{pap1} and ($ii$) the frictional interaction of the tube with its surroundings limit the residence time of flux tubes in the overshoot region?  We approach the problem by considering the nonlinearity of friction-induced instability (Sect.~\\ref{sec:friction}), including a discussion of the effects of finite perturbations on the magnetic buoyancy instability (Sect.~\\ref{ssec:delta}), and determine the ranges of flow parameters that allow the storage of sub-critical flux tubes subject to radial flows (Sect.~\\ref{sec:flows}).  The analyses and simulations were carried out in the framework of the thin flux tube approximation \\citep{spruit81}, which at present is the only existing approach that can deal with the small magnetically induced variations in density, pressure, and temperature corresponding to the high plasma $\\beta$ ($>10^5$) of the deep solar convection zone. ",
        "conclusions": " \\begin{itemize} \\item The flow instability driven by the frictional coupling of transversal MHD waves with external flows is nonlinear in the sense that the growth rate is a function of the initial perturbation amplitude. This is consistent with the results of the linear stability analysis of paper~III, in which the parameter $\\alpha$ is proportional to the speed of the perpendicular flow and thus to the amplitude of the subsequent perturbation. Therefore, the perturbation amplitude determines the $\\alpha$ parameter in the linear analysis.  \\item Significant buoyancy variations ($\\delta$-effect) along magnetic flux tubes can lead to a nonlinear buoyancy instability. The $\\delta$-effect most likely occurs in regions of large radial gradients of superadiabaticity such as the bottom of the overshoot layer, where a sufficient $\\delta$ variation along the tube ($\\Delta\\delta\\gtrsim 10^{-6}$) is easier to attain than in the upper layers.  \\item For flux tubes in the solar overshoot layer, we have established links between the perpendicular flow velocity amplitude, its spatial (azimuthal) extension, and the resulting radial displacement and instabilities induced by $\\delta$-effect or longitudinal flows.  \\item To store magnetic flux with flux densities between $10^4$-$10^5$~G for times comparable to the dynamo amplification time in the convective overshoot layer, the average flow speed and the flow duration must not exceed about 50~m~s~$^{-1}$ and 100~days, respectively, and that the azimuthal extension of the flow is lower than about $10^\\circ$. If we assume that these conditions prevail in the solar overshoot region, then magnetic fluxes of up to $10^{22}$~Mx can be stored within thin flux tubes during the dynamo amplification phase. \\end{itemize}"
    },
    "0910/0910.3296_arXiv.txt": {
        "abstract": "}[2]{{\\footnotesize\\begin{center}ABSTRACT\\end{center} \\vspace{1mm}\\par#1\\par \\noindent {~}{\\it #2}}}  \\newcommand{\\TabCap}[2]{\\begin{center}\\parbox[t]{#1}{\\begin{center} \\small {\\spaceskip 2pt plus 1pt minus 1pt T a b l e} \\refstepcounter{table}\\thetable \\\\[2mm] \\footnotesize #2 \\end{center}}\\end{center}}   \\newcommand{\\TableSep}[2]{\\begin{table}[p]\\vspace{#1} \\TabCap{#2}\\end{table}}  \\newcommand{\\FigCap}[1]{\\footnotesize\\par\\noindent Fig.\\  % \\refstepcounter{figure}\\thefigure. #1\\par}  \\newcommand{\\TableFont}{\\footnotesize} \\newcommand{\\TableFontIt}{\\ttit} \\newcommand{\\SetTableFont}[1]{\\renewcommand{\\TableFont}{#1}}  \\newcommand{\\MakeTable}[4]{\\begin{table}[htb]\\TabCap{#2}{#3} \\begin{center} \\TableFont \\begin{tabular}{#1} #4 \\end{tabular}\\end{center}\\end{table}}  \\newcommand{\\MakeTableSep}[4]{\\begin{table}[p]\\TabCap{#2}{#3} \\begin{center} \\TableFont \\begin{tabular}{#1} #4 \\end{tabular}\\end{center}\\end{table}}  \\newenvironment{references}% { \\footnotesize \\frenchspacing \\renewcommand{\\thesection}{} \\renewcommand{\\in}{{\\rm in }} \\renewcommand{\\AA}{Astron.\\ Astrophys.} \\newcommand{\\AAS}{Astron.~Astrophys.~Suppl.~Ser.} \\newcommand{\\ApJ}{Astrophys.\\ J.} \\newcommand{\\ApJS}{Astrophys.\\ J.~Suppl.~Ser.} \\newcommand{\\ApJL}{Astrophys.\\ J.~Letters} \\newcommand{\\AJ}{Astron.\\ J.} \\newcommand{\\IBVS}{IBVS} \\newcommand{\\PASP}{P.A.S.P.} \\newcommand{\\Acta}{Acta Astron.} \\newcommand{\\MNRAS}{MNRAS} \\renewcommand{\\and}{{\\rm and }} { We report time-series photometry for 16 variable stars located in the central part of the globular cluster NGC~6752. The sample includes 13 newly identified objects. The precision of our differential photometry ranges from 1~mmag at $V=14.0$ to 10~mmag at $V=18.0$. We detected 4 low amplitude variables located on the extended horizontal branch (EHB) of the cluster. They are candidate binary stars harboring sdB subdwarfs. A candidate  degenerate binary was detected about 2 mag below the faint end of the EHB. The star is blue and its light curve is modulated with a period of 0.47~d. We argue that some of the identified variable red/blue stragglers are ellipsoidal binaries harboring degenerate stars. They have low amplitude sine-like light curves and periods from a few hours to a few days. Spectroscopic observations of such objects may lead to  the detection of detached inactive binaries harboring stellar mass black holes or neutron stars. No binaries of this kind are known so far in globular clusters although their existence is expected based on the common occurrence of accreting LMXBs and millisecond pulsars. An eclipsing SB1 type binary was identified on the upper main sequence of the cluster. We detected variability of optical counterparts to two X-ray sources located in the core region of NGC~6752. The already known cataclysmic variable B1=CX4 experienced a dwarf nova type outburst. The light curve of an optical counterpart to the X-ray source CX19 exhibited modulation with a period  0.113~d. The same periodicy was detected in the HST-ACS data. The variable is located on the upper main sequence of the cluster. It is an excellent candidate for a close degenerate binary observed in quiescence. }  {\\bf Key words:} {\\it stars: dwarf novae - novae, cataclysmic variables -- globular clusters: individual: NGC 6752 -- stars: horizontal-branch -- binaries: eclipsing} ",
        "introduction": "\\label{sect:intro} NGC~6752 is a nearby globular cluster whose low reddening and relatively high Galactic latitude ($(m-M)_{V}=13.02$, $E(B-V)=0.04$, $b=-25.6$~deg; Harris 1996) make it an attractive target for detailed studies. Optical counterparts for 12 out of 19 faint X-ray sources detected with $Chandra$ were reported by Pooley et al. (2002). This sample included ten likely cataclysmic variables (CVs) and 1-3 RS~CVn or BY Dra stars. Two of the candidate CVs were detected and studied earlier by Bailyn et al. (1996) based on HST/WFPC2 data. So far there have been no reports of dwarf novae type outbursts in the field of NGC~6752. The cluster is known to host five millisecond pulsars (D'Amico et al. 2002), one of which has  an optical counterpart (Ferraro et al. 2003; Bassa et al. 2003). A wide-field CCD based survey of NGC~6752 conducted by Thomson et al. (1999) led to the detection of eleven photometric variables, seven of which were classified as contact binaries, and three as SX~Phe stars. The issue of membership status of these variables remains open. In particular,  Rucinski (2000) argues that the group of contact binaries is dominated by field interlopers.  The cluster has a rich population of blue horizontal branch (BHB) and extreme horizontal branch (EHB) stars (Buonanno et al. 1986; Momany et al, 2002). Until recently not a single photometric variable was known among the horizontal branch stars in NGC~6752. Catelan et al. (2008) reported possible detection of a BHB pulsator, while Kaluzny \\& Thompson (2008) detected four variables located on the BHB/EHB of the cluster. Only one spectroscopic binary was found among a few dozen BHB/EHB stars observed by Moni Bidin et al. (2008). This  result was unexpected, as close binaries are common among field sdB stars (Maxted et al. 2001; Napiwotzki et al. 2004).  The photometric survey presented here was conducted as a part of the CASE project (Kaluzny et al. 2005). It is complementary to the wide-field study presented by Thompson et al. (1999). The new data, collected with a larger telescope at a finer xpatial scale, allowed a more detailed study of the central part of the cluster. Since the pioneering work by Mateo et al. (1990) several globular clusters have been surveyed with CCD photometry for faint variables, and in particular for main sequence binaries. These surveys, usually conducted with 1-m class telescopes, are very incomplete in the cluster core regions, arising mainly from the saturation of stellar profiles of densely packed bright stars. This saturation is hard to avoid as exposures times have to be at least a few minutes long to assure sufficient S/N for the relatively faint main-sequence stars. The problem can be partially overcome by using larger telescopes with finer pixel scales producing images with a reduced number of saturated stars. An even better solution is to use imaging capabilities of the $HST$, but so far only one cluster has been systematically surveyed for variability with this instrument (47 Tuc, Gilliland et al. 2000; Albrow et al. 2001). Results published for a few other clusters observed with the $HST$ are based on fragmentary data.  It is expected  that most of the eclipsing variables are located in  the central parts of globular clusters. First, by definition, roughly 50\\% of all stars are observed within the half-light radius of a given cluster. Second, dynamical considerations indicate that most of the binary stars migrate toward  the core regions due to mass segregation (Stodolkiewicz 1986; Hut 1992). As we show below, our adopted observing strategy along with careful  data reduction has allowed a successful search for variability in  a large sample of stars in the central part of NGC~6752. Good quality photometry extends down to $V\\approx 21$~mag and the light curves of the brightest stars show an rms of about  1~mmag. ",
        "conclusions": "\\label{sect:summary}  We have obtained time series $BV$ photometry for about 40 thousand stars from the central area of the globular cluster NGC~6752. The observing strategy and careful reduction of the data resulted in a photometric precision  of 4  mmag stars with $V<16.0$  and 1-2 mmag at $V<15$, rising to 0.01 mag at $V\\approx 18.5$. The  sample included 72 hot subdwarf candidates. No pulsation variability was detected for any of them. We have detected, however, four low-amplitude variables which may be binary EHB/BHB stars. Spectroscopic follow up is needed to reveal their actual nature. Preliminary periods were established for two of them. A faint blue variable with $V\\approx 20.7$ was also detected. Its light curve is likely periodic with $P\\approx 0.47$~d. The variable is a good candidate for a degenerate binary belonging to the cluster. One of four periodically variable blue stragglers is a likely $\\gamma$~Dor star. If confirmed, it would be the first variable of this type detected in globular clusters.  We detected  three variable yellow/red stragglers with $15.0<V<16.3$. Two of them  show periodic sine-like modulation of their light curves. Another periodic variable was detected at $V=18.4$ on the right side of the upper-main sequence of the cluster. We propose that their light-curve modulation is related to binarity, and  that they are good candidates for ellipsoidal variables hosting degenerate components. Our hypothesis is based on two main arguments. First, it is known that field X-ray novae (binaries hosting stellar mass black holes) spend most of their time in quiescence, showing only an ellipsoidal variability (Remillard \\& McClintock 2006). Second, the known optical counterparts to millisecond pulsars located in GCs usually occupy positions either to the red or to the blue of the main sequence on the H-R diagram (Ferraro et al. 2001, 2003). As GCs contain rich population of active X-ray binaries, one may expect that they also harbor a large number of non-accreting, degenerate binaries. Unambiguous detection of non-accreting degenerate binaries in GCs would open a new, exciting field of research, and may lead to the detection of the first known stellar-mass black holes in GCs. The list of possible quiescent degenerate binaries in GCs is limited to a few faint X-ray sources of unknown nature (Heinke et al. 2003). The binary nature of red/yellow stragglers from NGC~6752 can be checked by obtaining a few medium resolution spectra per object. These spectra  could be  used to estimate the mass function for confirmed binaries, further constraining the nature of possible compact components.  We have detected a likely dwarf nova type outburst of one of the CV candidates located in the  cluster core region. Periodic variability was detected for the optical counterpart of the faint X-ray source CX19 located in the core region of the cluster. The star is located very close to the upper main sequence of the cluster, and is unlikely to be an ordinary cataclysmic variable.  We propose that the object it is an  excellent candidate for a close degenerate binary caught in  quiescence. Despite being located in the crowded area the variable is accessible for ground based spectroscopy.       \\Acknow{  Research of JK is supported by the grant MISTRZ from the Foundation for the Polish Science and by the grant N N203 379936 from the Ministry of Science and Higher Education. I.B.T. acknowledges the support of NSF grant AST-0507325. Based on observations made with the NASA/ESA Hubble Space Telescope, and obtained from the Hubble Legacy Archive, which is a collaboration between the Space Telescope Science Institute (STScI/NASA), the Space Telescope European Coordinating Facility (ST-ECF/ESA) and the Canadian Astronomy Data Centre (CADC/NRC/CSA). }"
    },
    "0910/0910.5365_arXiv.txt": {
        "abstract": "{We present the results from a combined study of the average X--ray spectral and timing properties of 14 nearby AGN.} {We investigate whether a ``spectral-timing\" AGN correlation exists, similar to the one observed in Cyg X-1, compare the two correlations, and constrain possible physical mechanisms responsible for the X--ray emission in compact, accreting objects.} {For 11 of the sources in the sample, we used all the available data from the \\rxte\\ archive, which were taken until the end of 2006.  There are 7795 \\rxte\\ observations in total for these AGN, obtained over a period of $\\sim 7-11$ years. We extracted their 3--20 keV spectra and fitted them with a simple power-law model, modified by the presence of a Gaussian line (at 6.4 keV) and cold absorption, when necessary.  We used the best-fit slopes to construct their sample distribution function, and we used the median of the distribution, and the mean of the best-fit slopes, which are above the 80th percentile of the distributions, to estimate the mean spectral slope of the objects. The latter estimate is more appropriate in the case when the energy spectra of the sources are significantly affected by absorption and/or reflection effects. We also used results from the literature to estimate the average spectral slope of the  three remaining objects.} {The AGN average spectral slopes are not correlated either with the black hole mass or the characteristic frequencies that were detected in the power spectra. They are  positively correlated, though, with the characteristic frequency when normalised to the sources black hole mass.  This correlation is similar to the spectral-timing  correlation that has been observed in Cyg X-1, but not the same.} {The AGN spectral-timing correlation can be explained if we assume that the accretion rate determines both the average spectral slope and the characteristic time scales in these  systems. The spectrum should steepen and the characteristic frequency should increase, proportionally, with increasing accretion rate. We also provide a quantitative expression between spectral slope and accretion rate. Thermal Comptonisation models are broadly consistent with our result, and can also explain the difference between the spectral-timing correlations in Cyg X-1 and AGN, but only if the ratio of the soft photons' luminosity to the power injected to the hot corona is proportionally related to the accretion rate.} ",
        "introduction": "The X--ray variability properties of AGN have been extensively studied during the past twenty years. Significant progress has been achieved in the estimation of their X--ray power spectral density functions (PSDs), which (among other things) can be helpful in the search for characteristic time scales  in these objects. This progress has been made possible  with the combined use of monitoring \\rxte\\ light curves (which are up to 5--10 years long in many cases) and shorter (1 to a few days long), high signal-to-noise, \\xmm\\ and {\\it Chandra} light curves. The results have shown that the PSD has a $\\sim -2$ power law shape at high frequencies and then,  below a characteristic ``break - frequency\", $\\nu_{\\rm bf}$, it flattens to a slope of $\\sim -1$ (e.g. Uttley, M$^{\\rm c}$Hardy, \\& Papadakis 2002; Markowitz \\etal\\ 2003, M$^{\\rm c}$Hardy \\etal\\ 2004). Uttley \\& M$^{\\rm c}$Hardy (2005) (UM05 hereafter) list $\\nu_{\\rm bf}$ estimates for 14 nearby AGN. Using these results, M$^{\\rm c}$Hardy \\etal\\ (2006) (M06 hereafter) demonstrated that the corresponding ``break timescale\", $T_{\\rm br}=1/\\nu_{\\rm bf}$,  increases with increasing black hole mass, M$_{\\rm BH}$, and for a given M$_{\\rm BH}$, it decreases with increasing accretion rate, $\\dot{m}_{\\rm E}$ (in units of the Eddington limit).  Knowledge of the X-ray properties of the Galactic black hole binaries (GBHs) has also advanced substantially during the past twenty years. Power spectral studies of Cyg X-1 in particular have been advanced considerably. Pottschmidt \\etal\\ (2003; hereafter P03) for example have used many \\rxte\\ observations between 1998 and 2001 to study the long-term evolution of the PSD. They also studied the energy spectrum of the source and presented convincing evidence that its timing and spectral properties are closely linked: the characteristic time scales become shorter  as the spectrum steepens. Axelsson, Borgonovo \\& Larsson (2006; hereafter A06), using several archival \\rxte\\ observations, which  cover all spectral states of the source, detected the same ``spectral-timing properties\" correlation as well. Shaposhnikov and Titarchuk (2006) also found that photon index and characteristic PSD frequencies are positively correlated in Cyg X-1.  The question whether AGN vary in a manner similar to that of GBHs is a long standing one. The availability of the better quality AGN light curves over the last few years (resulted from intense and extended RXTE monitoring campaigns to sample variability on very long time scales) has allowed a more quantitative comparison between AGN, Cyg X-1 and other GBHs (GRS1915+105 for example,  in M06). In this work we used archival \\rxte\\ observations of 11 AGN, together with data from the literature for 3 more objects, to estimate their average spectral slope and compare it with the characteristic frequencies that have been detected in their power spectra. We show that a positive  ``average spectral slope - characteristic frequencies\" correlation exists, and  we argue that this correlation is driven by accretion rate, for a given black hole (BH) mass: objects with a  higher accretion rate relative to the Eddington limit should also have a steeper spectrum and a shorter characteristic time scale.  We also compared the spectral-timing relation we found in AGN with a similar relation that has been detected in one GBH binary, namely Cyg X-1. The two relations differ, but by an amount that can be explained if we take  into account the BH mass difference in AGN and Cyg X-1. Our results support the idea that  AGN and GBHs vary in the same way.  They also have interesting implications regarding the nature of the X--ray source in AGN and GBHs. ",
        "conclusions": "We have used 7795 \\rxte\\ observations of 11 AGN, obtained over a period of $\\sim 7-11$ years, to extract their 3--20 keV spectra. We fitted  them with a simple ``power-law + Gaussian line\" model, and we used the best-fit slopes to construct their sample distribution function. We used the median of the distributions, and the mean of the best-fit slopes which are above the 80th percentile of the distributions, to estimate the mean spectral slope of the objects (the latter estimate is more appropriate in the case when the energy spectra of the sources are significantly affected by absorption and/or reflection effects). We also used results from literature to estimate the average spectral slope of three more objects. The fourteen AGN that we consider in this work are nearby, X--ray bright objects, whose X--ray light curves have been studied in the past. Their PSDs have been accurately estimated, and characteristic ``break frequencies\" have been detected in them.  When we combine the mean spectral slope estimates with the $\\nu_{\\rm br}$ estimates listed in UM05, we find that: {\\it objects with steeper mean energy spectra have shorter characteristic time scales as well}. This is the first time that such a ``spectral -- timing\" correlation is detected in AGN. The results we reported in Section 4.3 suggest that this spectral-timing relation in AGN can be parametrised as follows: $\\Gamma\\approx 1.3\\nu_{\\rm norm}^{0.07}$.  \\subsection{The spectral-timing correlation in AGN}  The easiest way to explain this correlation is to assume that, for an AGN with a given BH mass, the accretion rate determines both  the PSD characteristic frequencies (this has already been shown by M06) and their energy spectral shape as well: the higher the accretion rate, the steeper the mean energy spectrum will be. Such a positive $\\Gamma-\\dot{\\rm m}_{\\rm E}$ relation in AGN has already been suggested/shown in the past (e.g. Porquet \\etal\\ 2004;  Bian 2005; Shemmer \\etal\\ 2006; Saez \\etal\\ 2008). A positive $\\Gamma-\\dot{\\rm m}_{\\rm E}$ relation has also been detected recently in 7 GBHs by Wu \\& Gu (2008), when they accrete at a rate higher than $\\sim 1\\%$ of the Eddington limit. Our results are consistent with the results reported in these papers.  Using the above results and the M06 results, we calculate that $\\nu_{\\rm norm}($Hz$\\times$M$_{\\odot})\\approx 3000\\dot{\\rm m}_{\\rm E}$. When we  replace $\\nu_{\\rm norm}$ in the relation of  $\\Gamma\\approx 1.3\\nu_{\\rm norm}^{0.07}$ we found in this work, we find that: $\\Gamma \\approx 2.3\\dot{\\rm m}_{\\rm E}^{0.07}$.  This formula between spectral slope and accretion rate could be helpful in deriving a rough estimate for the  accretion rate of the  distant AGN that are detected in deep X-ray surveys, assuming that their variability properties are similar to the properties of nearby AGN (this is supported by recent studies, see e.g. Papadakis \\etal\\ 2008).  In the context of thermal Comptonisation models, $\\dot{\\rm m}_{\\rm E}$ can affect the spectral slope $\\Gamma$ as it controls the strength of the soft disc photons, hence the cooling of the thermal  plasma in the X-ray emitting corona. The greater the cooling by seed photons incident on the plasma, the softer the resulting X-ray power-law spectra are. Indeed, we found that Comptonisation models are consistent with the $\\Gamma - \\dot{\\rm m}_{\\rm E}$  relation that our results imply,  but  only if $L_{\\rm s}/L_{\\rm diss}\\propto \\dot{\\rm m}_{\\rm E}$.  \\subsection{The comparison with Cyg X-1}  Both P03 and A06 have used data from \\rxte\\ observations that lasted for $\\sim 1-2$ ks, and were performed every $\\sim 5-10$ days over a period of many years. If ``time scales\" scale with BH mass in accreting objects, then $1-2$ ksec in Cyg X-1 should correspond to a period of at least $\\sim 3-6$ years in objects with M$_{\\rm BH}> 10^6$ times the mass of the black hole in Cyg X-1  (like most AGN in our sample). This suggests that each one of the AGN points in Fig.~3 corresponds to just one of the Cyg X-1 points plotted in the same figure. We found that the AGN and Cyg X-1 ``$\\Gamma - \\nu_{\\rm norm}$\" relations are similar but not the same. In Section 5 we discussed briefly some implications of this result. In the paragraphs below we discuss the implications of our results in more detail.  The main aim of the  discussion below is to investigate the constrains that the AGN spectral-timing relation, and its comparison with the similar relation in Cyg X-1, impose on thermal Comptonisation models, based on the particular assumption that both the spectral and timing properties of accreting systems are driven by accretion rates variations. We point out that other interpretations are also possible; see for example Kylafis \\etal\\ (2008) for an alternative explanation of the Cyg X-1 spectral-timing relation, which does not assume that X--rays are produced by thermal Comptonisation. We also point out that, given the small number of objects in our sample, and the unavoidable uncertainty in the derived parameters of the AGN spectral-timing relation, the values of the various parameters in the equations below are somehow uncertain.  As we showed in Section 6.1, the AGN spectral-timing relation and the  M06 results imply that $\\Gamma_{\\rm AGN} \\approx 2.3  \\dot{m}_{\\rm E}^{0.07}$.  If X--rays in AGN  are produced by thermal Comptonisation, we expect\\footnote{The discussion in these paragraphs is based on the assumption of an outflowing corona, hence we adopt the results of Malzac \\etal, 2001. Similar conclusions can also be drawn if we assume a static corona.} that $ \\Gamma_{\\rm AGN} = 2.2 (L_{\\rm s}/L_{\\rm diss})^{0.07}$.  Therefore, the observations [i.e. $\\Gamma_{\\rm AGN} \\approx 2.3  \\dot{m}_{\\rm E}^{0.07}$] are consistent with the thermal Comptonisation model predictions [$ \\Gamma_{\\rm AGN} = 2.2 (L_{\\rm s}/L_{\\rm diss})^{0.07}$], only if  \\begin{equation} (L_{\\rm s}/L_{\\rm diss})\\approx 2 \\dot{m}_{\\rm E}. \\end{equation}  \\noindent An obvious implication of this result is that, if a certain fraction, say $f_{\\rm s}$, of the total accretion power, $P_{\\rm tot}$ $(=\\eta\\dot{m}_{\\rm E}c^2,$ where $\\eta$  is the efficiency of the accretion process), is converted to disc luminosity (i.e. $L_{\\rm s}=f_{\\rm s}P_{\\rm tot}$), while  $L_{\\rm diss}=f_{\\rm diss}P_{\\rm tot}$, then the ratio $(f_{\\rm s}/f_{\\rm diss})$ should not remain constant for a given object, but should rather increase with increasing accretion rate.  Suppose that X--rays from Cyg X-1 in its low/hard state (LH) are also produced by thermal Comptonisation. In this case,  thermal Comptonisation models predict that $\\Gamma_{\\rm CygX-1} = 2.2 (L_{\\rm s}/L_{\\rm diss})^{0.13}$, or  \\begin{equation} \\Gamma_{\\rm CygX-1}\\approx  2.4 \\dot{m}_{\\rm E}^{0.13}, \\end{equation}  \\noindent if we accept that equation (1) holds in this case as well. According to K\\\"{o}rding \\etal\\ (2007), the normalisation of the $\\nu_{\\rm norm}$ vs $\\dot{m}_{\\rm E}$ relation for the hard-state GBHs is $\\sim 8$ times smaller than the normalisation in the case of the AGN in our sample. So, if  $\\nu_{\\rm norm, Cyg X-1/LS}\\approx 375\\dot{m}_{\\rm E}$, as opposed to $\\nu_{\\rm norm, AGN}\\approx 3000\\dot{m}_{\\rm E}$, then  $\\dot{m}_{\\rm E}\\approx 3\\times 10^{-3}\\nu_{\\rm norm, Cyg X-1/LS}.$ We can now substitute  $\\dot{m}_{\\rm E}$ in equation (2) to determine the $\\Gamma -\\nu_{\\rm norm}$ relation for Cyg X-1 in LH state:  \\begin{equation} \\Gamma_{\\rm CygX-1/LS}\\approx  1.1 \\nu_{\\rm norm}^{0.13}. \\end{equation}  \\noindent The dashed line in the top panel of  Fig.~3 indicates this relation. The agreement between the {\\it predicted} $\\Gamma-\\nu_{\\rm norm}$ relation  and the Cyg X-1 data is rather good up to $\\nu_{\\rm norm}\\sim 70$. The discussion so far suggests the following picture:  a) X-rays from the AGN we studied are produced by thermal Comptonisation.  The $\\Gamma-\\nu_{\\rm norm}$ relation we observe is consistent with the predictions of thermal Comptonisation models but only if the $(L_{\\rm s}/L_{\\rm diss})$ ratio, and hence the $(f_{\\rm s}/f_{\\rm diss})$ ratio as well, increase proportionally with accretion rate.  b) X-rays in Cyg X-1 are also produced by thermal Comptonisation. Taking into account the fact that the normalisation of the $\\nu_{\\rm norm}-\\dot{m}_{\\rm E}$ relation is $\\sim 8$ times smaller in Cyg X-1 than in the AGN in our sample (K\\\"ording \\etal\\ 2007), the {\\it predicted} $\\Gamma-\\nu_{\\rm norm}$ relation agrees well with the Cyg X-1 data up to $\\nu_{\\rm norm}\\approx 70$.  c) In the case when $\\Gamma_{\\rm AGN}=\\Gamma_{\\rm Cyg X-1}$, we expect that ($L_{\\rm s}/L_{\\rm diss})_{\\rm Cyg X-1}^{0.13}=(L_{\\rm s}/L_{\\rm diss})_{\\rm AGN}^{0.07}$, and $\\dot{m}_{\\rm AGN,E}^{0.07}=\\dot{m}_{\\rm Cyg X-1,E}^{0.13}$ (using equation 1). Consequently, $\\dot{m}_{\\rm E,AGN}=\\dot{m}_{\\rm E,Cyg X-1}^{1.9}$ and, since  $\\dot{m}_{\\rm E}<1$, AGN should operate on a lower accretion rate than Cyg X-1 when the spectral slope is the same in both systems. Furthermore, since $\\Gamma_{\\rm AGN}\\approx 1.3 \\nu_{\\rm norm,AGN}^{0.07}$  and $\\Gamma_{\\rm Cyg X-1}\\approx 1.1 \\nu_{\\rm norm,Cyg X-1}^{0.13}$ (when $\\Gamma\\la 2.1-2.2$), then  $\\Gamma_{\\rm AGN}=\\Gamma_{\\rm Cyg X-1}$ implies that $1.3 \\nu_{\\rm norm,AGN}^{0.07}=1.1 \\nu_{\\rm norm,Cyg X-1}^{0.13}$, and $\\nu_{\\rm norm,Cyg X-1}\\approx 3.5\\nu_{\\rm norm,AGN}^{0.55}$. Therefore, when the spectral slope is the same (and less than $\\sim 2.1-2.2$) in AGN and Cyg X-1, the former should operate at a lower accretion rate but their characteristic time scales should be shorter than those in Cyg X-1 (when normalised to the respective BH mass), because the normalisation of the  AGN $\\dot{m}-\\nu_{\\rm norm}$ relation is significantly larger than the normalisation of the respective Cyg X-1 relation in LH state.  The vertical dot-dashed line in the top panel of Fig.~3  indicates the value $\\nu_{\\rm norm}=70$.  For Cyg X-1,  this normalised frequency corresponds to $\\nu_{1}=1$ Hz and $\\nu_2=5$  Hz for the centroid frequency of the  low and higher frequency Lorentzians, respectively. As $\\nu_{\\rm norm}$ (i.e. the accretion rate) increases even more, then i)  the ratio $\\nu_{1}/\\nu_2$ increases as well (see Fig.~2 in A06), ii) the contribution of the Lorentzians to the root mean square variability amplitude decreases, and iii) the ``bending\" power-law component in the PSD appears and its contribution to the  source variability amplitude increases (see Fig.~7 in A06), i.e. the Cyg X-1 power spectrum changes from a hard to a soft-state shape. At $\\nu_{\\rm norm}\\approx 70$, the average {\\it observed} spectral slope of Cyg X-1 is $\\sim 2.1$ (see Fig.~3), while  at higher frequencies the spectral slope is steeper. Consequently, the region defined by  $\\Gamma<2.1$ and $\\nu_ {\\rm norm}<70$ corresponds to the ``hard state region\" for Cyg X-1 in the ``spectral-timing plane\" shown in Fig.~3.  The discrepancy between the predicted spectral-timing relation and the Cyg X-1 data when $\\Gamma\\ma 2.1$ and $\\nu_{\\rm norm}\\ma 70$ cannot be explained by the fact that the normalisation of the  $\\nu_{\\rm norm} - \\dot{m}_{\\rm E}$ relation increases by a factor of $\\sim 8$ in the high/soft state.  If that were the case,  the  spectral-timing relation in this state should be similar to the one defined by equation (3), but with a {\\it smaller} normalisation (opposite to what we observe).  One possibility is that the  $(L_{\\rm s}/L_{\\rm diss})-\\dot{m}_{E}$ relation (equation 1) changes in the high/soft state. However, in this case we would have to accept that the X--ray source does  {\\it not}  operate in the same way in AGN and Cyg X-1: when  $\\nu_{\\rm norm}>70$, both Cyg X-1 and the AGN in our sample follow the same $\\nu_{\\rm norm} - \\dot{m}_{\\rm E}$ relation (M06, K\\\"ording \\etal\\ 2007). Therefore, as long as $\\nu_{\\rm norm}>70$, a given  $\\nu_{\\rm norm}$ value implies the  same accretion rate in both systems. The fact that the AGN spectral-timing relation is valid up to $\\nu_{\\rm norm}\\approx 1000$  should then imply that, for the same accretion rate,  the  ``$(L_{\\rm s}/L_{\\rm diss}) - $ accretion rate\" relation is different in AGN and Cyg X-1.  In Cyg X-1, $\\Gamma\\approx 2.1$  implies that $2.2 (L_{\\rm s}/L_{\\rm diss})^{0.13}\\approx 2.1$, and hence  $(L_{\\rm s}/L_{\\rm diss})\\approx 0.7$. Another possible explanation then for the discrepancy between the Cyg X-1 data and the predicted spectral-timing relation above $\\nu_{\\rm norm}=70$ is the following: equation (1) holds until $(L_{\\rm s}/L_{\\rm diss})\\approx 0.7$, at which point the hot corona is significantly cooled down, and the thermal X--ray emission component is weak. It is possible then that at high accretion rates  a separate, possibly non-thermal, X--ray component emerges, and dominates the X-ray emission in the soft state. If that is the case, the $\\Gamma- L_{\\rm s}/L_{\\rm diss}$ relation we have assumed above is not valid, hence the predicted spectral-timing relation does not fit the data any more.  Even if the picture drawn above is correct, there are important issues regarding the relation between AGN and GBH states which still remain unresolved.  In particular, the answer to the question whether the AGN in our sample are ``soft\" or ``hard\" state systems is far from clear. There are indications that they are analogous to Cyg X-1 in its soft state. For example, the radio emission in Cyg X-1 in its LH state is enhanced. On the other hand, Panessa \\etal\\ (2007) have shown that, for the same $L_{\\rm X}/L_{\\rm E}$ ratio, the radio luminosity of  Seyfert galaxies is $\\sim 8-10$ times lower than the radio luminosity of hard-state GBHs, even when the BH mass difference is properly taken into account (7 of the objects in our sample are also included in their sample). Furthermore, the AGN in our sample follow the soft-state ``characteristic time scale -- accretion relation\" in GBHs, and  a ``bending\" power-law is the dominant component in the power spectrum of those objects with good enough light curves to accurately study their PSD (e.g. NGC~4051, M$^{\\rm c}$Hardy \\etal\\ 2004; NGC~3227 and  NGC~5506, UM05; MCG~6-30-15, M$^{\\rm c}$Hardy \\etal\\ 2005;  and perhaps NGC~3783, Summons  \\etal\\ 2007), implying a soft-state like PSD in these objects. High quality light curves for low accretion rate AGN are necessary to investigate whether any AGN with hard-state like power spectra exist or not.  However, although the radio emission strength and the timing properties of many objects in our sample are soft-state like, their spectral properties are {\\it not}, as the average spectral slope is smaller than $2.1$ in most cases. If indeed the {\\it spectral} hard-to-soft state transition corresponds to $(L_{\\rm s}/L_{\\rm diss})\\sim 0.7$, at which point the thermal corona emission is weak, and a different component dominates the X--ray emission, then we should expect this transition to happen when $\\Gamma\\ma 2.15$ in AGN. Given the AGN $\\Gamma-\\nu_{\\rm norm}$ relation, this slope corresponds to $\\nu_{\\rm norm}\\approx 2500$. At even higher normalised  frequencies, we should then expect the AGN spectral-timing relation to break (like in Cyg X-1 at $\\nu_{\\rm norm}\\approx 70$). Obviously, more data are necessary to confirm that this is indeed the case in AGN.  The data so far suggest that while in AGN the {\\it timing} properties transition from hard-to-soft state happens at least as low as  $\\nu_{\\rm norm}\\approx 100$, the {\\it spectral} properties transition  should happen at much higher accretion rates. This is opposite to what we observe in Cyg X-1, where both the spectral and timing properties change from the hard to soft state, at the same accretion rate (indicated by the value $\\nu_{\\rm norm}\\approx 70)$. Perhaps the timing properties in accreting objects are determined by accretion disc variations only, and the hard-to-soft state transition materialises at a certain accretion rate, irrespective of whether the soft disc luminosity is strong enough for $(L_{\\rm s}/L_{\\rm diss})> 0.7$, i.e. strong enough to cool down the hot corona. In this case, we would expect the AGN hard-to-soft timing properties transition to appear at $\\nu_{\\rm norm}\\approx 70$ as well. Due to the cooler disc temperature though, the AGN spectral soft-to-hard state transition happens at higher accretion rates (i.e. at a higher $\\nu_{\\rm norm}$ value in the spectral-timing plane of Fig.~3). Further progress in understanding the relation between AGN and GBHs can be made when we know how the accretion rate determines the characteristic frequency in accreting compact objects (assuming that it is just the accretion rate that determines $\\nu_{\\rm norm}$ in these systems)."
    },
    "0910/0910.2545_arXiv.txt": {
        "abstract": "Blue stragglers have been found in all populations. These objects are important in both stellar evolution and stellar population synthesis. Much evidence shows that blue stragglers are relevant to primordial binaries. Here, we summarize the links of binary evolution and blue stragglers, describe the characteristics of blue stragglers from different binary evolutionary channels, and show their consequences for binary population synthesis, such as for the integrated spectral energy distribution, the colour-magnitude diagram, the specific frequency, and the influences on colours etc.. ",
        "introduction": "Blue stragglers (BSs) are an important population component in stellar evolution as well as in star clusters. These objects have remained on the main sequence for a time exceeding that expected from standard stellar evolution theory, and they may affect the integrated spectra of their host clusters by contributing excess spectral energy in the blue and UV bands. Many mechanisms, including single star models and binary models, have been presented to account for the existence of BSs (see the review of \\cite{str93}). At present, it is widely believed that more than one mechanism plays a role for the produce of BSs in one cluster and that binaries are important or even dominant for the production of BSs in open clusters and in the field (\\cite{lan07,dal08,sol08}). Binaries may produce BSs by way of mass transfer, coalescence of the two components, binary-binary collision and binary-single star collision. The collision of binary-binary or binary-single may lead binaries to be tighter or farther apart, and it is relevant to dynamics and environment in the host cluster. In this contribution, we are only concerned with BSs resulting from the evolutionary effect of primordial binaries, i.e. mass transfer and coalescence of two components. ",
        "conclusions": "In this contribution, we summarized the linkage of binary evolution and BSs, and BS characteristics. As well, we showed binary population synthesis results of BSs from primordial binary evolution, such as the distribution on CMD, the contribution to ISED, the specific frequency and the influences on colours. This work was in part supported by the Chinese National Science Foundation (Grant Nos. 10603013 and 10973036,10821061 and 2007CB815406) and Yunnan National Science Foundation (Grant No. 08YJ041001)."
    },
    "0910/0910.0589_arXiv.txt": {
        "abstract": "We show that an object classified as a galaxy in on-line data bases and revealed on sky survey images as a distant ring galaxy is a rare case of polar ring galaxy where the ring is only slightly inclined to the equatorial plane of the central body. Imaging information from the Sloan Digital Sky Survey (SDSS) indicates that the diameter of the ring is about 36 kpc. The SDSS data was combined with long-slit spectroscopic observations and with  Fabry-P\\'{e}rot Interferometer H$\\beta$ mapping obtained at the Russian Academy of Sciences 6-m telescope. We derived the complex morphologies of this presumed ring galaxy from a combination of SDSS images and from the kinematical behaviour of the central body and of the ring, and determined the stellar population compositions of the two components from the SDSS colours, from the spectroscopy, and from models of evolutionary stellar synthesis. % The metallicity of the ring material is slightly under-abundant. The total luminosity and the total mass of the system are not extreme, but the rather high M/L$\\simeq$20 indicates the presence of large amounts of dark matter.  We propose two alternative explanations of the appearance of this object. One is a ring formed by two semi-circular and tight spiral arms at the end of a central bar. The apparent inclination between the ring and the central body, and a strange kink at the North-East end of the ring, could be the result of a warp or of precession of the ring material. The object could, therefore, be explained as an extreme SBa(R) galaxy. The other possibility is that we observe a Polar Ring Galaxy where the inner object is an S0 and the ring is significantly more luminous than the central object. The compound object would then be similar to the NGC 4650A galaxy, but then it would be a rare object, with a polar component only modestly inclined to the equatorial plane of the central body. Arguments for (and against) both explanations are given and discussed, with the second alternative being more acceptable. ",
        "introduction": "\\label{txt:Obs_and_Red}  \\subsection{SDSS data} \\label{txt:sdss}   The SDSS \\citep[][]{York2000,DR5} imaging data were used to provide morphological parameters for the galaxy, colours of the central object, and colours and colour distributions for the ring. Since the source extraction routines of SDSS are not optimized for the detection and measurement of very faint surface brightness extended objects, SDSS provides only the photometry for the central component and also for those ``shreds'' of ring identified as ``stars'' by the SDSS detection software. Therefore, specialized photometry in the SDSS bands was applied to measure the properties of the central component and of the ring using the low surface brightness detection methods developed and described in \\citet{LSB_phot}.  Specifically, a deep image of the galaxy and of the ring was created by combining the $g$, $r$, and $i$ images with weights, as explained in \\citet{LSB_phot}. While no physical meaning can easily be attached to photometry using such a wide and ill-defined photometric band, the resulting image is significantly deeper than the individual SDSS images and is useful to investigate the object morphology. The deep $gri$ combined image is shown in Figure~\\ref{Fig:RG1_aper}. % An outermost isophote was used to trace automatically an irregularly-shaped aperture that collects all the light from even very faint surface brightness parts of the galaxy. %  At this point, a few remarks about the appearance of RG1 are in order. Figure~\\ref{Fig:RG1_aper} shows a relatively bright, elliptical-shaped central body encircled by an $\\sim$25-arcsec elliptical ring. Two roundish and relatively compact non-stellar objects are visible near the ring; one is to the East (galaxy E) and the other is in the opposite side (galaxy W), slightly more distant from the centre and fainter than galaxy E. These bodies are reminiscent of the proposed ``impactor'' galaxy in a scenario for the  formation of collisional ring galaxies. In addition, the ring shows a loop-like extension or kink to the North-East.  In order to perform integral and surface brightness photometry we used concentric circular annuli centered on the galaxy's photo-centre \\citep{LSB_phot}, with radii increasing outward by one SDSS pixel. The average surface brightness and the average colour were determined in each annulus.   \\subsection{Long-slit spectroscopy with the 6m BTA telescope  of the RAS } \\label{txt:long_slit_obs}  First spectroscopic data were obtained % with the 6-meter Large Altazimuthal Telescope (BTA) of the Special Astrophysical Observatory of Russian Academy of Science (SAO RAS) on February 16/17 2001 (see Table~\\ref{t:Obs} for details). The 2001 observations used the Long-Slit spectrograph UAGS (Afanasiev et al. 1995) at the telescope prime focus equipped with a Photometrics CCD detector with 1024$\\times$1024 pixels each $24\\times24\\, \\mu m^2$ in size. Two 130\\arcsec\\ long-slit spectra were obtained in blue and in red with a grating having 651 grooves/mm yielding a dispersion of 2.4~\\AA/pixel and a spectral resolution of 6--7~\\AA\\,. The slit position (see Fig.~\\ref{Fig:RG1_aper}) was chosen to cross the central bright region close to the direction of the major axis of the inclined ring. The wavelength ranges of the obtained spectra are given in Table~\\ref{t:Obs}. A 2\\arcsec\\ wide slit was used in both cases. The scale along the slit was 0.39\\arcsec/pixel, very similar to that of the SDSS images. Reference spectra of an Ar--Ne--He lamp were recorded before or after each observation to provide  wavelength calibration. Spectrophotometric standard stars from Bohlin (1996) were observed for flux calibration.  New spectral observations with the BTA were obtained by AM on January 15/16 and on February 12/13, 2008, as also detailed in Table~\\ref{t:Obs}. These were collected using the long-slit mode of the SCORPIO universal focal reducer (Zasov et al.  2008). A volume-phase holographic grating with 1200 grooves mm$^{-1}$ was used, providing a dispersion of 0.84 \\AA\\,pixel$^{-1}$ and yielding spectra that covered the range $\\sim$5600\\AA\\, to $\\sim$7310\\AA. The spectral resolution of these spectra was $\\sim$5 \\AA\\, and the spatial sampling was 0.36 arcsec pixel$^{-1}$. The locations of the entrance slit for the two observed positions are shown in Figure~\\ref{Fig:RG1_aper}.  The primary data reduction included cosmic-ray removal with the {\\tt MIDAS}\\footnote{{\\tt MIDAS} is an acronym for the European Southern Observatory package --- Munich Image Data Analysis System.} software package, and bias subtraction and flat-field correction with the {\\tt IRAF}\\footnote{{\\tt IRAF}: the Image Reduction and Analysis Facility is distributed by the National Optical Astronomy Observatories, which is operated by the Association of Universities for Research in Astronomy, In. (AURA) under cooperative agreement with the National Science Foundation (NSF).} software package. For subsequent reduction of the long-slit spectra we used {\\tt IRAF}. In case of UAGS data, after wavelength mapping and night sky subtraction, each 2D frame was corrected for atmospheric extinction and was flux calibrated. 2D spectra of different spectral regions were combined after that stage.  RG1 has an SDSS spectrum in the public archives. Since this was obtained with a single fiber, it samples only the central three arcsec of the central body. We extracted the innermost three arcsec spectrum from our long-slit reduced UAGS 2D spectrum and compared it with the SDSS spectrum. Although the comparison shows that the two spectra match within 10\\%, we assumed that the SDSS spectrum has a more accurate flux distribution and constructed the correction function based on it.  The reason for this assumption is that SDSS is observing simultaneously with 640 fibers, where some are targeted on flux standards. This implies that each SDSS spectroscopic measurement has a few standard stars out of the 640 spectra collected and their fluxes are used for calibration. For 6m blue and red parts of the final spectrum where observed separately in not-so-perfect weather conditions. Additionally, blue and red parts of the flux standards were observed separately. Thus, from our point of view, ``many and simultaneous'' calibrations are better than ``few and separate'' and we selected the SDSS spectrum to derive the correction function. The derived function was applied to the reduced UAGS 2D spectrum and the result is shown in Figure~\\ref{Fig:2D-spectrum} together with the 1D spectra of the central component and the outer part of the ring.   To obtain the line-of-sight velocity distribution along the slit we used a classical cross-correlation method described in detail by \\citet{Zasov00}. All emission lines were measured with the MIDAS programs described in detail in \\citet{SHOC,Sextans}.    \\begin{figure*} {\\centering \\includegraphics[clip=,angle=-90,width=17.0cm]{fig2a.eps} \\includegraphics[clip=,angle=-90,width=8.0cm]{fig2b.eps} \\includegraphics[clip=,angle=-90,width=8.0cm]{fig2c.eps} } \\caption{{\\it Upper panel:} Part of the reduced  BTA UAGS 2D spectrum of RG1 for $\\rm PA \\sim 58^{\\circ}$. NE is up. The two arcsec wide slit was positioned on the central component and along the major axis. The ring region exhibits the redshifted [\\ion{S}{ii}], H$\\alpha$, [\\ion{N}{ii}],  [\\ion{O}{iii}], H$\\beta$, H$\\gamma$, and [\\ion{O}{ii}] emission lines with measurable intensities on a visible, though faint, continuum that seems stronger on the blue side of the image. The two Ca II H \\& K lines and the CH G-band absorption are visible near the blue end of the central object spectrum. The spectrum of the central object and of both sides of the ring is easily visible for a distance of $\\pm$14\\arcsec\\ along the slit.  At the adopted distance of 250 Mpc 1\\arcsec\\ = 1.2 kpc and the image vertical extent is $\\sim$33.6 kpc. Horizontal lines to the right of the panel show parts of 2D spectrum, which were averaged and used in Section~\\ref{txt:abund}. {\\it Lower panel:} The 1D spectra extracted from the 2D spectrum observed at $\\rm PA \\sim 58^{\\circ}$: the central component of RG1 (left, region {\\it c}) and outer part of the ring (right, region {\\it 2}). Some detected emission and absorption lines have been marked. The spectrum of the central component is equivalent with the one in the SDSS database. \\label{Fig:2D-spectrum}} \\end{figure*}   \\subsection{Observations with the scanning Fabry-P\\'{e}rot interferometer}  The UAGS spectra showed that the H$\\beta$ line from the ring is fairly well-defined and not confused with other spectral features. It is therefore beneficial to attempt a two-dimensional kinematic mapping with field spectroscopy.  Observations were performed during the night of 26-27 October 2008 with the scanning Fabry--P\\'{e}rot Interferometer (FPI) mounted within the SCORPIO focal reducer (Afanasiev \\& Moiseev 2005) at the prime focus of the SAO RAS 6-m telescope. The desired spectral interval in the neighbourhood of the redshifted H$\\beta$ line was selected by a narrow-band filter with full-width at half maximum (FWHM) of 21\\AA\\,. The free spectral interval between the neighbouring interference orders was 18\\,\\AA\\, ($\\sim$1100 km/s). The FPI resolution, defined as the $FWHM$ of the instrumental profile, was $1.5$\\AA\\, ($\\sim90$ km/s) for a  0.55\\AA\\, sampling. The detector was a $2048\\times2048$ EEV 42-40 CCD operating in on-chip binned $4\\times4$ pixel mode to reduce readout time and match the resultant pixel to the seeing size. The final image scale and field of view were 0\\farcs7/pixel and 6\\farcm1$\\times$6\\farcm1, respectively.  We collected 32 successive interferograms of the object with different spacings of the  FPI plates. The total exposure was 96~min and the seeing during the observations varied from  1\\farcs5 to 2\\farcs0.  The observations were reduced  using an IDL-based software package described by Moiseev (2002) and Moiseev  \\& Egorov (2008). Following the primary reduction, night-sky line subtraction, and wavelength calibration, the frames were combined into a data cube where each pixel in the $512\\times512$ field contains a 32-channel spectrum centered on the H$\\beta$ line sampled at 0.55\\AA\\ per bin. The  final angular resolution after optimal smoothing matches the  2\\farcs4 seeing.  \\begin{figure*} \\centerline{ \\includegraphics[width=9cm]{Fig_3a.eps} \\includegraphics[width=9cm]{Fig_3b.eps}} \\caption{Results of FPI observations: intensity distribution of the H$\\beta$ line emission (left panel, with a linear intensity scale) and line-of-sight velocity field of the ionized gas (right panel). The coordinate origin coincides with the dynamical centre of the object.} \\label{fig:FPIimage} \\end{figure*}  Figure~\\ref{fig:FPIimage} shows an image of the object in the H$\\beta$ line (left panel) obtained by Gaussian fitting the line emission in the FPI data cube, and subtracting a properly scaled combination of the line-free channels that contain the continuum contribution. The right panel of Figure~\\ref{fig:FPIimage} shows the distribution of radial velocities over the same image. Note that there is line emission in the central object and in the bridge connecting it to the ring. Note also the enhancements in the ring itself; these features shall be discussed below. Figure~\\ref{fig:FPIimage} shows also an emission patch $\\sim$10\\arcsec\\ away from the ring at (--$\\Delta\\alpha$, $\\Delta\\delta$)=(13, 13); this corresponds to a faint patch in the SDSS $gri$ deep image (see Figure~\\ref{Fig:RG1_aper}).% ",
        "conclusions": "\\label{txt:summ}  We presented above SDSS photometry, long-slit spectroscopy, and FPI data about a possible ring galaxy. The observations have shown that:  \\begin{enumerate} \\item The central body shows a S\\'{e}rsic surface brightness distribution with $n\\simeq1.2$, is much redder than the ring, and exhibits emission lines on top of an absorption spectrum.  \\item The ring is blue, with strong emission lines, maintains the same colour through its width, and is connected to the CB by a linear structure. The H${\\alpha}$-to-H${\\beta}$ emission line ratio indicates A$_B$=2.7 in the ring. This, and the forbidden line ratios, indicate that the ring metallicity is 12+log(O/H)$\\sim$8.32, lower than expected considering the ring luminosity.  \\item The CB stars have a velocity dispersion commensurate with the maximal ring rotation velocity. Their rotation velocity at the outermost locations joins smoothly with the ring rotation.  \\item The ionized gas in the ring shows an increasing rotation curve almost to its outermost point. The systemic velocity of each ring segment decreases with galacto-centric distance. The plane of the ring is tilted by 58$^{\\circ}\\pm10^{\\circ}$, or 73$^{\\circ}\\pm11^{\\circ}$, from the plane of the CB.  \\item The rotation curve indicates a total dynamical mass of 4$\\times 10^{11}$ M$_{\\odot}$ within the outer edge of the ring. Since the entire object has M$_g$=--20.35, assuming  a solar absolute magnitude of 5.45 mag in the same spectral band implies a total object luminosity of $\\sim2\\times 10^{10}$ L$_{\\odot}$. Therefore RG1 has an overall M/L$\\simeq$20, rather high for galaxies and implying significant amounts of dark matter in the system. \\end{enumerate}  Sil'chenko \\& Moiseev (2006) showed that NGC 7742, an Sb galaxy with a nuclear star-forming ring and lacking a bar, has two global exponential stellar disks with different scale lengths and a circumnuclear inclined gaseous disk with a radius of 300 pc. The galaxy has a global gaseous counterrotating subsystem, which they interpret as a result of a past minor merger, including the appearance of the nuclear star-forming ring lacking a global bar; the ring might be produced as a resonance feature by tidally-induced oval distortions of the global stellar disk.  The connection between the CB and the ring in RG1, and the smooth match of the rotation curves of the CB (from stellar absorption lines) to that of the ring (from the emission lines), indicates that RG1 could probably be an extreme case of barred spiral where the two arms form a complete circle. However, the kinematics described here, in particular the second solution obtained for the inclination, would indicate an extreme warp of the disk; while the inner regions are reasonably flat, from about 25 arcsec out to the end of the luminous body the disk turns by 60--70$^{\\circ}$ within a few arcsec.  It is also possible that this is another case of polar-ring galaxy, like the well-known NGC4650A with a relatively  small central stellar body and massive self-gravitating stellar and gaseous polar ring. The polar disk could have a spiral structure induced by the non-axisymmetric potential of CB; this is similar to that of a bar in a disk galaxy but without being a real bar, matching the simulations of Theis et al. (2006).  The similarities with respect to NGC 4650A can be summarized as (a) a CB with a near-exponential profile, with a S\\'{e}rsic index $n\\simeq1$, (b) a bright and presumably massive, self-gravitating ring or disk, and (c) a spiral or quasi-spiral structure in the polar disk/ring.  We reconsider here the issue of metal abundance in light of the possible nature of the objet. We remind the reader that the metal abundance measurement shown in Figure~\\ref{fig:line_ratios} showed a fairly constant abundance of 12+log(O/H)=8.30 to 8.65. The measurements covered ring radii 4$\\leq$R$\\leq$15 arcsec; given that the 25 mag arcsec$^{-1}$ surface brightness level is reached at R$\\simeq$12 arcsec, this implies that our abundance measurements for the ring cover normalized radii 0.4$\\leq \\frac{R}{R_{25}} \\leq$1.25. We compare this with the metallicity gradients measured by van Zee et al. (1998) for a sample of spiral galaxies, where their figure 12 shows considerable oxygen gradients over this range.  In contrast to the situation regarding spiral galaxies, as shown by van Zee et al. (1998) as well as by Denicolo et al. (2001), very few PRGs have been studied to the determine their metal abundances. A few examples are UGC 5600 (Shalyapina et al. 2002) where no gradient was found, NGC 7468 (Shalyapina et al. 2004) with an inverted gradient where the outer parts are more metal-rich than the inner parts, and AM1934-563 (Brosch et al. 2007) with a normal gradient for the galaxy but a constant metallicity for the ring.  We conclude that this finding for RG1 vs. a fairly steady gradient for more metal-poor regions further from the centre for disk galaxies and a variety of gradients for PRGs, argues in favour interpreting the extended feature in RG1 as indeed a ring, and not part of a disk.  Another argument in favour of considering RG1 a PRG and not a disky spiral is the character of the kinematics of the gas vs. stars in the central region. Not only have we observed counter-rotation of these two components, but the motions there occur on different planes. This is different for spiral galaxies that show counter-rotation; the motions there take place in the same plane and while counter-rotation is seen along the major axis, only a zero velocity gradient is seen along the minor axis (e.g., Sil'chenko et al. 2009).  We summarize the arguments for or against these two possibilities for the nature of RG1 in the Truth Table~\\ref{t:truth}. The option of a PRG has more + signs than the other option thus, on the balance of probabilities, we accept it over the alternative of a barred spiral with tight spiral arms and a strong disk warp. Therefore RG1 could be a rare case of a galaxy where the polar ring shows only a small inclination with respect to the equatorial plane of the CB.  \\begin{table} \\caption{Two options to explain RG1} \\label{t:truth} \\begin{tabular}{ccc} \\\\ \\hline Property        & Disk & PRG   \\\\ \\hline Elongated CB connected with ring  & + & -- \\\\ Relative brightness ring/CB & -- & +  \\\\ Colours CB/ring & + & + \\\\ Ring metal abundance & -- & + \\\\ Ring radial brightness distribution & -- & + \\\\ Central gas (H$\\beta$) vs. stars kinematics & -- & + \\\\ CB velocity dispersion & -- & + \\\\ Tully-Fisher & + & -- \\\\ Cluster neighbourhood & + & -- \\\\ \\hline \\end{tabular} \\end{table}"
    },
    "0910/0910.0540_arXiv.txt": {
        "abstract": "Using the IRAC images from the {\\it Spitzer c2d} program, we have made a survey of mid-infrared outflows in the $\\rho$ Ophiuchi molecular cloud. Extended objects that have prominent emission in IRAC channel 2 (4.5 \\micron) compared to IRAC channel 1 (3.6 \\micron) and stand out as green objects in the three-color images (3.6 \\micron \\ in blue, 4.5 \\micron \\ in green, 8.0 \\micron \\ in red) are identified as mid-infrared outflows. As a result, we detected 13 new outflows in the $\\rho$ Ophiuchi molecular cloud that have not been previously observed in optical or near-infrared. In addition, at the positions of previously observed HH objects or near-infrared emission, we detected 31 mid-infrared outflows, among which seven correspond to previously observed HH objects and 30 to near-infrared emission. Most of the mid-infrared outflows detected in the $\\rho$ Ophiuchi cloud are concentrated in the L1688 dense core region. In combination with the survey results for Young Stellar Objects (YSOs) and millimeter and sub-millimeter sources, the distribution of mid-infrared outflows in the $\\rho$ Ophiuchi molecular complex hints a propagation of star formation in the cloud in the direction from the northwest to the southeast as suggested by previous studies of the region. ",
        "introduction": "Mass outflows play an essential role in the process of star formation and have been found to be ubiquitous in various stages of star formation \\citep{arc07,bal07}. The specific angular momentum of a star-forming molecular core is about four orders of magnitude times higher than that of a typical T Tauri star \\citep{bod95}. One way to transfer the excessive angular momentum from the star-forming cores so that the circumstellar material can be accreted onto the central stars is through mass outflows \\citep{sha07}. Mass outflows from YSOs have been observed at different wavelengths. In the visual, high velocity jets and HH objects with a typical velocity of 100 - 300 km s$^{-1}$ trace material ejected by the star or shocked surrounding medium \\citep{rei01}. Near-infrared molecular hydrogen emission features, which usually have velocities of several tens km s$^{-1}$, trace H$_2$ gas in the jets or in the entrained surrounding medium \\citep{eis00}. In the millimeter wavelength, high velocity CO gas with velocity in the range from a few to ten km s$^{-1}$ have been observed \\citep{wu04}. The high velocity CO gas probes the swept or entrained medium. The different velocity ranges and excitation conditions for HH objects, near-infrared molecular hydrogen emission, and CO outflows indicate the transfer of energy and momentum from the high velocity material to the surrounding medium.  Shocked gas in outflows has abundant atomic and molecular hydrogen line emission in the mid-infrared. Hydrodynamic simulations of outflows from YSOs by \\citet{sr05} have shown that outflows have strong H$_2$ line emission in all the IRAC four bands. In fact, mid-infrared observations with IRAC aboard of {\\it Spitzer} have discovered a lot of new outflows \\citep{mer08} and plenty new features of previously known outflows \\citep{nor04,smith06,tei08}, showing that {\\it Spitzer} IRAC imaging is a powerful tool in the survey of outflows, particular for those that are deeply embedded in clouds due to the much less extinction in the mid-infrared compared to the optical and near-infrared wavelengths. The {\\it Spitzer} legacy program {\\it c2d} performed a complete imaging of five nearby star-forming regions with the IRAC and MIPS instruments \\citep{eva03,eva09}, which provides a valuable data set to search for mid-infrared outflows in these regions. In this paper we present the results of a mid-infrared outflow survey of the $\\rho$ Ophiuchi molecular cloud based on the {\\it c2d} archive images.  The $\\rho$ Ophiuchi molecular cloud is one of the closest star-forming regions at a distance of about 130 pc \\citep{wil08}. It is intermediate in star formation activity compared to the isolated star formation in the Taurus cloud and the rich clusters in the Orion cloud. The large scale structure of the $\\rho$ Ophiuchi cloud has been revealed by extensive molecular $^{13}$CO line mapping \\citep{lor89a,lor89b}. Millimeter and submillimeter surveys of the central part of the cloud have revealed the clumpy structure of the cloud and have found about 55 starless cores with masses in the range of 0.02 to 6.3 M$_\\sun$ \\citep{motte98,johnstone00,you06}. The cloud has been surveyed for young stars at wavelengths from X-ray \\citep{montmerle83,imanishi01,ozawa05}, visual \\citep{wilking87,wilking05}, near-infrared \\citep{gre92,allen02}, mid- and far-infrared \\citep{you86,wilking01,pad08}, to millimeter and submillimeter \\citep{sta06,andrews07}. Based on these surveys, \\citet{wil08} compiled a list of 316 young stars in L1688, the main cloud of the $\\rho$ Ophiuchi complex. About 46 HH objects \\citep{wilking97,wu02,phe04}, 119 H$_2$ emission features \\citep{gro01,gom03,kha04}, 16 high velocity CO outflows \\citep{bontemps96,bussmann07} have been observed in the $\\rho$ Ophiuchi complex. For a summary of the $\\rho$ Ophiuchi complex we refer to the recent review by \\citet{wil08} ",
        "conclusions": ""
    },
    "0910/0910.5223_arXiv.txt": {
        "abstract": "We present Hubble Space Telescope optical coronagraphic polarization imaging observations of the dusty debris disk HD 61005.  The scattered light intensity image and polarization structure reveal a highly inclined disk with a clear asymmetric, swept back component, suggestive of significant interaction with the ambient interstellar medium.  The combination of our new data with the published 1.1 $\\mu$m discovery image shows that the grains are blue scattering with no strong color gradient as a function of radius, implying predominantly sub-micron sized grains.  We investigate possible explanations that could account for the observed swept back, asymmetric morphology.  Previous work has suggested that HD 61005 may be interacting with a cold, unusually dense interstellar cloud. However, limits on the intervening interstellar gas column density from an optical spectrum of HD 61005 in the Na I D lines render this possibility unlikely. Instead, HD 61005 may be embedded in a more typical warm, low-density cloud that introduces secular perturbations to dust grain orbits.  This mechanism can significantly distort the ensemble disk structure within a typical cloud crossing time.  For a counterintuitive relative flow direction---parallel to the disk midplane---we find that the structures generated by these distortions can very roughly approximate the HD 61005 morphology.  Future observational studies constraining the direction of the relative interstellar medium flow will thus provide an important constraint for future modeling. Independent of the interpretation for HD 61005, we expect that interstellar gas drag likely plays a role in producing asymmetries observed in other debris disk systems, such as HD 15115 and $\\delta$ Velorum. ",
        "introduction": "Nearly two dozen dusty debris disks surrounding nearby stars have now been spatially resolved at one or more wavelengths.  Many of these systems show clear similarities.  For example, the radial architecture of several debris disks can be understood in terms of a unified model of steady-state dust production via collisions in a parent planetesimal belt (e.g., Strubbe \\& Chiang 2006).  However, while the observed structure of many systems is ring-like \\citep{Kalas06,Wyatt08}, most disks show substructure such as clumps, warps, offsets, and brightness asymmetries not explained in traditional steady-state collisional grinding models.  These unexpected features have triggered a great deal of recent theoretical work.  The effects of massive planetesimal collisions, sandblasting by interstellar grains, close stellar flybys, dust avalanches, and secular and resonant perturbations by exoplanets have all been invoked to explain the observations (e.g., Moro-Martin et al. 2007, and references therein).  However, as many of these theories produce similar structures, further observational constraints are needed to better understand the key forces affecting disk structure and the circumstances in which they apply.  At a heliocentric distance of 34.5 pc \\citep{Perryman97}, the debris disk surrounding HD 61005 (SpT: G8 V; Gray et al. 2006), is a promising target for advancing our understanding in this area.  The significant infrared excess for this source ($L_{\\rm IR}/L_{*} = 2 \\times 10^{-3}$) was recently discovered as part of the Spitzer FEPS survey \\citep{Carpenter09}, indicating 60 K blackbody-emitting grains $\\gtrsim$ 16 AU from the star.  Follow-up Hubble Space Telescope (HST) coronagraphic imaging observations with the Near Infrared Camera and Multi-Object Spectrometer (NICMOS; HST/GO program 10527; D. Hines, PI) resolved the source (Hines et al. 2007, hereafter H07), revealing an unprecedented swept, asymmetric morphology, suggestive of significant interaction with the interstellar medium (ISM).  H07 suggested that this system could be a highly inclined debris disk, undergoing ram pressure stripping by the ambient ISM.  However, this interpretation requires an unusually high interstellar density for the low-density Local Bubble in which HD 61005 resides.  Furthermore, the single wavelength intensity image was insufficient to provide strong constraints on the dominant size of the scattering grains and the overall scattering geometry.  To further quantify the physical properties of grains seen in scattered light and the overall geometry of the system, we obtained optical coronagraphic polarimetry imaging observations of HD 61005 with the Advanced Camera for Surveys (ACS) onboard HST.  As demonstrated by \\citet{Graham07} for the case of AU Mic, polarization observations in scattered light are invaluable for breaking degeneracies between grain scattering properties and their spatial distribution.  Furthermore, the ACS data represent a factor of two improvement in angular resolution compared to the 1.1 $\\mu$m discovery observations.  In addition to these new imaging data, we also obtained a high resolution optical spectrum to characterize ambient interstellar gas surrounding this system.  In \\S 2, we describe the steps taken in observing and reducing these data.  In \\S 3, we discuss the results of these observations, their consequences for the system scattering geometry, and the additional constraints they provide when combined with the 1.1 $\\mu$m NICMOS image.  In \\S 4, we explore whether interactions with ambient interstellar gas can plausibly explain the observed swept, asymmetric morphology in this system.  We discuss the implications for these potential explanations in \\S 5 and summarize our findings in \\S 6. ",
        "conclusions": "\\subsection{Interstellar Gas and the HD 61005 Morphology}  The previous section explored whether disk/gas interaction can plausibly explain the unusual HD 61005 morphology.  Of the four scenarios considered, three are implausible, given the limits on the ambient interstellar gas density imposed by our optical spectrum.  The fourth scenario, secular perturbations from low density gas, is an attractive alternative, as this mechanism can significantly distort grain orbits well within a cloud crossing time.  Furthermore, the densities required by this scenario are typical of local interstellar clouds, which occupy up to $\\sim20$\\% of the local ISM.  Nevertheless, our preliminary modeling of this effect (Figures \\ref{vr} $-$ \\ref {vz}) can only produce disk morphologies in very rough agreement with the observations, suggesting that either additional physics needs to be incorporated into the current models, or that an altogether distinct physical mechanism is at work.  Indeed, the current models are simplistic, and their applicability is limited by several key assumptions:  \\begin{enumerate} \\item{{\\it Astrosphere Sizes:} As discussed by \\citet{Scherer00}, the toy models presented in Figure \\ref{vr} $-$ \\ref {vz} require that the disturbed grains be inside the astrospheric termination shock, such that the interstellar gas density and velocity can be approximated as constant. For the case of HD 61005, the termination shock distance is unknown. Furthermore, as HD 61005 is farther away than any star for which a direct astrospheric detection has been made \\citep{Wood04}, the astrosphere size may be difficult to constrain observationally.  In general, termination shock distances vary greatly, depending on the ambient ISM and stellar wind conditions (e.g., densities, temperatures, velocities, stellar activity).  For the case of the Sun only, the hydrodynamic models of \\citet{Muller06} show that the termination shock distance could easily vary between $\\sim 10$ AU and $500$ AU.  Observational astrosphere measurements of solar-type stars are consistent with these predictions\\footnote{ On a broader scale, no astrosphere detections have been made for stars earlier than G-type.  As a result, the typical effect of astrospheres on IS gas drag in debris disks surrounding A-type stars is difficult to reliably assess.} (Mann et al. 2006, and references therein). } \\item{{\\it Initial Conditions:} As noted in \\S 4.2.2 and Appendix B, our adopted initial orbital elements for the HD 61005 disk grains prior to the interstellar cloud interaction are highly uncertain, given a lack of information for the grain properties and underlying planetesimal population that collisionally replenishes the observed dust disk. Future long-wavelength observations sensitive to larger grains may be able to place tighter constraints on the distribution of sub-micron grains prior to the interstellar cloud interaction, as the distribution of large grains would likely reflect that of the parent bodies for the sub-micron population.  Furthermore, larger grains would not be significantly affected by interstellar gas drag on the same timescale as the sub-micron size grains traced in these observations.  The numerical techniques employed here and in \\citet{Scherer00} can be easily revised to accommodate an arbitrary initial disk architecture, provided the orbits are not highly eccentric, such that averaging over one orbit and applying Gauss' method is invalid.} \\item{{\\it Internal Disk Collisions:} In Appendix B, we estimate to order-of-magnitude that the collision time for submicron grains at $\\sim$70 AU is $\\sim$5000 yr. This collision time is somewhat longer than the timescale over which our model relaxes to a steady state --- essentially the time for gas drag to unbind a grain --- given in Appendix B as $\\sim$3000 yr. That the times are comparable supports the assumption of our models that each grain removed by gas drag is collisionally replenished. At the same time, the comparison of timescales underscores a shortcoming of our model --- that removal of grains by collisions is ignored. In reality, submicron grains should be removed from the system not only by gas drag, but also by collisions, in roughly equal proportions.  We defer to future work a comprehensive study that includes removal by collisions via a collisional cascade.} \\item{{\\it Planetary Configurations:} Finally, as illustrated by \\citet{Scherer00} (e.g., his Figure 2), the incorporation of planetary orbits can appreciably change the perturbed orbital elements from the case in which only IS gas drag is considered. This caveat is especially important for massive grains, or grains in close proximity to planetary orbits.  As a result, the models presented here should be treated with some caution if applied to typical planetary system scales ($\\lesssim 50$ AU; e.g., Kenyon \\& Bromley 2004).  This cautionary point may be particularly important for the case of HD 61005, as the origin of the brightness asymmetry between the northeast and southwest disk lobes (\\S 3.1.1) is unknown.  The agreement between the northeast and southwest deflected component position angles (\\S 3.1.3 and Figures \\ref{acs_image_log} $-$ \\ref{acs_image_lin}) suggests this asymmetry may originate from a physical mechanism entirely distinct from ISM interaction.  If the asymmetry is due to a massive perturber, the disk morphologies produced in Figures \\ref{vr} $-$ \\ref{vz} are likely to be inapplicable.  Resolved long wavelength observations sensitive to massive grains are needed to further explore this possibility.} \\end{enumerate}  In addition to the above uncertainties, a remaining ambiguity important for future IS gas drag modeling is the velocity of the putative cloud responsible for the HD 61005 morphology.  H07 noted that the star's tangential space motion is perpendicular to the disk midplane, in agreement with the relative flow vector suggested by initial inspection of the observed images.  However, while assigning the relative flow direction to the star's tangential velocity is appealing, the three scenarios explored in \\S 4 which assume this flow direction were found to be untenable.  Furthermore, velocities of local warm clouds can be comparable to the observed space motion of HD 61005 (\\S 4.2.2).  As a result, the star's tangential motion is not a reliable indicator of the cloud-star relative velocity.  Future spectroscopic observations may be able to detect the cloud directly (e.g., HST/GO Program 11674; H. Maness, PI), thereby providing key constraints on the ambient ISM density and velocity.  Such observations will greatly inform future modeling, as the preliminary interstellar gas drag models presented here suggest a counterintuitive relative motion parallel to the disk midplane, rather than perpendicular to it.  Though HD 61005 is difficult to assign to any known interstellar clouds (\\S 4.2.2), its galactic coordinates could plausibly associate it with either the G cloud or the Blue cloud.  The space velocities of both of these clouds suggest a relative motion dominated by the radial velocity component.  Therefore if HD 61005 is associated with either of these clouds, the relative motion is inconsistent with all models posited in \\S 4.  \\subsection{General Applicability of Interstellar Gas Drag}  The normalcy of the interstellar densities, velocities, and cloud sizes required by the secular perturbation model in \\S 4.2.2 suggests that IS gas drag can be important beyond HD 61005 in shaping debris disk morphologies.  Taking the simple models of \\S 4.2.2 and Appendix B at face value, several of the general morphological features produced in Figures \\ref{vr} $-$ \\ref{vz} are consistent with observed disk structures.  For example, the extreme brightness asymmetry in HD 15115 \\citep{Kalas07} may potentially result from interstellar gas drag, though a range of alternative explanations could explain this system as well (e.g., see list in \\S 1).  The bow structures seen in some of the face-on models in Figure \\ref{vr} are also reminiscent of the mid-infrared morphology observed around the A star, $\\delta$ Velorum, which was recently modeled as a purely interstellar dust phenomenon \\citep{Gaspar08}.  Finally, the middle panels in Figure \\ref{vrvz} show that warps, similar to that seen in $\\beta$ Pic (e.g., Mouillet et al. 1997), can in principle be produced for a relatively wide range of flow directions.  However, while IS gas drag can in principle produce commonly observed disk features, the rate at which gas drag is expected to affect the observations remains unclear. Beyond the uncertainties in the model physics described in \\S 5.1, the characteristics of warm, low density clouds are currently uncertain, as detailed knowledge of them is limited to clouds residing predominantly within 15 pc of the Sun \\citep{Redfield08}.  As a result, our understanding of typical cloud sizes, shapes, and total volumetric filling factor remains rudimentary.  A key finding in this area, however, is that a significant fraction of nearby warm clouds appear to exhibit filamentary morphologies, which would limit the average interaction time between a given disk and cloud, likely reducing the rate at which IS gas perturbations produce an observable effect.  This concern is particularly important for the case of disks surrounding early-type stars, as grains traced in scattered light tend to be larger in disks surrounding A-type stars than in their later type counterparts, owing to the larger radiation pressure blowout size.  As such, the scattered-light morphologies for A-star disks require correspondingly longer cloud-disk interaction times to be noticeably affected.  The timescale for a given grain to become unbound under IS gas drag increases approximately as the square root of the grain size \\citep{Scherer00}.   \\subsection{Interstellar Grains and the HD 61005 Morphology}  Finally, we note that all models posited in \\S 4 consider only the role of interstellar gas, ignoring the potential effects of interstellar grains. \\citet{Artymowicz97} investigated IS sandblasting of debris disks surrounding A stars and found that sandblasting has a negligible effect on the observed structure, as radiation pressure blows out most incoming interstellar grains before they are allowed to intersect the disk.  However, under this framework, only grains with $\\beta \\geq 1$ are ejected.  Thus Figure \\ref{radpr} shows that radiation pressure does not protect the HD 61005 disk, as it does in A-stars.  Nevertheless, even if radiation pressure does not protect the disk against sandblasting, the stellar wind might, as only large interstellar grains with sizes greater than a few $\\times 0.1 \\mu$m are allowed to enter astrospheres freely without deflection \\citep{Linde00}.  Thus it is likely that interstellar sandblasting can only plausibly compete with interstellar gas drag if the astrosphere is smaller than or comparable to the observed debris disk size \\citep{Mann06}.  The size of the HD 61005 astrosphere is unconstrained by present observations.  In general, observations and models of astrospheres surrounding solar-type stars show sizes in the range $\\sim 10-10^3$ AU, depending on the ambient ISM and stellar wind conditions (see discussion in \\S 5.1).  Thus with a characteristic disk size of $\\lesssim 70$ AU, it is not clear whether typical interstellar grains can intersect the HD 61005 disk.  Detailed modeling of sandblasting is outside the scope of this paper. However, future theoretical work should investigate the effects of sandblasting on debris disks surrounding solar-type stars. Calculations of the ISM density required for sandblasting to eject an observable flux of grains, the disk morphologies produced in this case, and the timescale for which sandblasting can be sustained would significantly aid in differentiating between this explanation and the gas drag models presented here."
    },
    "0910/0910.5824_arXiv.txt": {
        "abstract": "We have used the zCOSMOS-bright 10k sample to identify 3244 {\\em Spitzer/MIPS} \\tfm-selected galaxies with $0.06< S_{24 \\, \\rm \\mu m}\\lsim 0.50 \\, \\rm mJy$ and $I_{\\rm AB}<22.5$, over 1.5~deg$^2$ of the COSMOS field, and studied different spectral properties, depending on redshift. At $0.2<z<0.3$, we found that different reddening laws of common use in the literature explain the dust extinction properties of $\\sim 80$\\% of our infrared (IR) sources, within the error bars. For up to 16\\% of objects, instead, the \\ha/\\hb ratios are too high for their IR/UV attenuations, which is probably a consequence of inhomogenous dust distributions.   In only a few of our galaxies at  $0.2<z<0.3$ the IR emission could be mainly produced by dust heated by old rather than young stars. Besides, the line ratios of $\\sim 22$\\% of our galaxies suggest that they might be  star-formation/nuclear-activity composite systems.   At $0.5<z<0.7$, we estimated galaxy metallicities for 301 galaxies: at least 12\\% of them are securely below the upper-branch mass-metallicity trend, which is consistent with the local relation. Finally, we performed a combined analysis of the $\\rm H_\\delta$ equivalent-width  versus $\\rm D_n(4000)$ diagram for 1722 faint and bright \\tfm galaxies at $0.6<z<1.0$, spanning two decades in mid-IR  luminosity.  We found that, while secondary bursts of star formation are necessary to explain the position of the most luminous IR galaxies in that diagram, quiescent, exponentially-declining star formation histories can well reproduce the spectral properties of $\\sim 40$\\% of the less luminous sources.  Our results suggest a transition in the possible modes of star formation at total IR luminosities  $L_{\\rm TIR} \\approx (3\\pm2) \\times 10^{11} \\, \\rm L_\\odot$. ",
        "introduction": "\\label{sec-intro}   Investigating the composition of the extragalactic IR background (Puget et al.~1996; Elbaz et al.~2002; Dole et al.~2006) at $z<1$ is of key importance to understand galaxy evolution and the global decline of the star formation rate density in the last 8 Gigayears (Gyr). At $z>1$, the IR background is mainly dominated by  luminous and ultra-luminous IR galaxies (LIRGs and ULIRGs, with total IR [$5-1000 \\, \\rm \\mu m$] luminosities $10^{11}<L_{\\rm TIR}<10^{12} \\, \\rm L_\\odot$ and  $L_{\\rm TIR}> 10^{12} \\, \\rm L_\\odot$, respectively). However, these populations are progressively less important at lower redshifts, becoming rare by the present epoch. The most typical IR sources at $z \\lsim 0.7$ are IR normal galaxies, i.e. those with total IR luminosities $L_{\\rm TIR}< 10^{11} \\, \\rm L_\\odot$. Their study is fundamental to clarify the production of the bulk of the IR background in the second half of cosmic time.    Our knowledge of IR galaxies at low redshifts was first possible thanks to studies conducted with the {\\em Infrared Astronomical Satellite (IRAS)} and the {\\em Infrared Space Observatory (ISO)} (see e.g. Aussel et al.~1999; Genzel \\& Cesarsky~2000, for a  review), which revealed that both star formation and nuclear activity were associated with the dust emission.  More recently, with the advent of the {\\em Spitzer Space Telescope} (Werner et al.~2004) and the {\\em Akari Telescope} (Matsuhara et al.~2006), a much deeper insight into the nature of the $z<1$ IR galaxy population has been possible. We have now characterised their luminosity evolution (e.g. Le Floc'h et al.~2005; Caputi et al.~2007; Magnelli et al.~2009), stellar masses (e.g. Caputi et al. 2006a,b), and spectral energy distributions (SEDs; e.g. Rowan-Robinson et al.~2005; Rocca-Volmerange et al.~2007; Takagi et al.~2007; Bavouzet et al.~2008; Kartaltepe et al.~2009; Symeonidis et al.~2009), among others (see Soifer et al.~2008, for a review).     Spectroscopic surveys at different wavelengths can add very valuable information that help understand the nature of the $z<1$ IR sources. In addition to providing accurate redshift determinations, spectroscopic data allow us to characterise other aspects of IR galaxies, such as their extinction properties, metallicities and general chemical composition.  The initial optical/near-IR spectroscopic studies of IR galaxies at  $z<1$ targeted {\\em IRAS} and {\\em ISO} sources (e.g. Kim et al.~1995; Veilleux et al.~1995, 1999; Franceschini et al.~2003; Flores et al.~2004; Arribas et al.~2008; Garc\\'{i}a-Mar\\'{i}n et al.~2009). In the {\\em Spitzer} era, and in the light of larger spectroscopic campaigns, more systematic studies of IR galaxies and their environments have been possible (e.g. Choi et al.~2006; Papovich et al.~2006; Caputi et al.~2008 [C08 hereafter], 2009; Desai et al.~2008).   In particular, optical spectra are useful to disentangle the power source of the IR emission, which  generally occurs when the galaxy dust re-processes the energy provided by the ultra-violet (UV) photons of young stars or active nuclei. However, in some cases, the observed IR emission might not be associated with any on-going stellar or nuclear activity, but be rather produced by dust heated by old stars (e.g. Gordon et al.~2000; Kong et al.~2004; Cortese et al.~2006; but also see Young et al.~2009), or even direct stellar emission. The importance of these mechanisms within the IR galaxy population is not yet well established. The identification of `passive' IR sources is important to quantify the degree of contamination in the IR-derived cosmic star formation rate densities. Spectroscopic data can be very helpful to elucidate the presence of star formation in sources that are presumably dominated by old stellar populations.  Spectral features also carry the imprints of the galaxy star formation history. Thus, for IR galaxies,  optical spectra can be used to investigate at what stage of the galaxy history the IR phase is produced.  Locally, the rare LIRGs and ULIRGs are known to be mainly the consequence of galaxy mergers that induce bursts of star formation and nuclear activity (e.g. Armus et al.~2007). Instead, more typical IR galaxies at $z=0$ are usually less disturbed, and appear to form their stars in a more regular way.  There are several clues that suggest that, in the past, when LIRGs and ULIRGs were more common,  both the burst-like and quiescent modes of star formation could also lead to the IR phase (e.g. Elbaz et al.~2007; Daddi et al.~2008). In spite of these different pieces of evidence, it is not yet clearly established whether there exists a transition in the possible modes of star formation between galaxies of different IR luminosities. One of the main aims of this work is to study such a transition in IR galaxies at $0.6<z<1.0$.   The Cosmic Evolution Survey (COSMOS; Scoville et al.~2007) has been designed to probe galaxy evolution, star formation and the effects of large-scale structure up to high redshifts. The COSMOS field is defined by its {\\em Hubble Space Telescope/Advanced Camera for Surveys (HST/ACS)}  2  deg$^2$ coverage (Koekemoer et al.~2007). There are multiple follow up observations carried out in the COSMOS field, ranging from X-rays to radio wavelengths. Among the photometric imaging observations, COSMOS includes the full and homogeneous coverage of the field with the {\\em Spitzer Infrared Array Camera} ({\\em Spitzer/IRAC}; Fazio et al.~2004) and the {\\em Multiband Photometer for Spitzer} ({\\em MIPS}; Rieke et al.~2004), as part of {\\em Spitzer} cycle-2 and 3 Legacy Programs (Sanders et al.~2007).   COSMOS also comprises a large spectroscopic follow-up program named zCOSMOS (Lilly et al.~2007), that is being performed with the {\\em Visible Multiobject Spectrograph} ({\\em VIMOS}; Le F\\`evre et al.~2003) on the {\\em Very Large Telescope (VLT)}. The zCOSMOS survey has two parts, one of which consists in the observation of an $I<$ 22.5 AB mag limited sample of 20,000 galaxies over 1.7 deg$^2$  of the COSMOS field  (zCOSMOS-bright).  Here, we make use of the first half of spectra observed and analysed in zCOSMOS-bright, which constitute the so-called `zCOSMOS-bright 10k sample' (Lilly et al.~2009).   This paper presents the second part of a study we have carried out to analyse the optical spectral properties of \\tfm-selected  galaxies in the COSMOS field, using the zCOSMOS-bright 10k sample.  In this second part, we analyse 3244  IR sources with $0.06 < S_{24 \\, \\rm \\mu m} \\lsim  0.50 \\, \\rm mJy$, which constitute the largest optical spectroscopic sample of mid-IR-selected galaxies analysed to date in the redshift range $0<z\\lsim 1$. The results presented in this paper complement our first study of the optical spectra of brighter {\\em MIPS} \\tfm sources (C08), and focus on the most outstanding aspects discussed in this previous work. The study of these new fainter galaxies is of particular importance:  they largely outnumber our previous sample and, especially,  span the typical IR luminosities that dominate the IR background at $z<1$ (Le Floc'h et al.~2005; Caputi et al.~2007; Magnelli et al.~2009), covering a range in $L_{\\rm TIR}$ of e.g. a few times $10^{10}$ to $9 \\times 10^{10} \\, \\rm L_\\odot$ at $z=0.5$, and between  $L_{\\rm TIR} \\approx 9 \\times 10^{10}$ and $5\\times 10^{11} \\, \\rm L_\\odot$ at $z=1$.   The layout of this paper is as follows: in Section \\S\\ref{sec_sample}, we describe the deep {\\em MIPS} \\tfm and zCOSMOS-bright datasets and the cross-correlation of the two catalogues. In Section \\S\\ref{sec_rflum}, we give details about the rest-frame mid-IR luminosities for our galaxies. We present our results in Sections \\S\\ref{sec_linelowz} and \\S\\ref{sec_linehighz},  based on a variety of line diagnostics in different redshift ranges: we extensively discuss the dust extinction properties of our galaxies at $0.2<z<0.3$, and study metallicities and stellar masses at $0.5<z<0.7$.  Later on, we perform a combined analysis of our present faint and previous brighter IR samples at $0.6<z<1.0$ to investigate the possible mechanisms of star formation in IR galaxies, and how they vary with IR luminosity.  Finally, we summarise and discuss our results in Section \\S\\ref{sec_concl}.  We adopt throughout a cosmology with $\\rm H_0=70 \\,{\\rm km \\, s^{-1} Mpc^{-1}}$, $\\rm \\Omega_M=0.3$ and $\\rm \\Omega_\\Lambda=0.7$.   All quoted magnitudes and colours refer to the AB system. ",
        "conclusions": "\\label{sec_concl}   In this paper, we have identified 3244 \\tfm-selected galaxies with $0.06< S_{24 \\, \\rm \\mu m}\\lsim 0.50 \\, \\rm mJy$  within the zCOSMOS-bright 10k sample, and studied different optical spectral properties depending on their redshifts. Our sample spans e.g. total IR luminosities $\\sim 3 \\times 10^{10}< L_{\\rm TIR} < 9 \\times 10^{10} \\, \\rm L_\\odot$ at $z=0.5$, and  $\\sim9 \\times 10^{10} < L_{\\rm TIR} < 5 \\times 10^{11} \\, \\rm L_\\odot$ at $z=1$. This means that our sources are IR normal galaxies at $z<0.5$, and a mixture of IR normal galaxies and LIRGs at $0.5<z<1.0$. This works complements our previous study of the zCOSMOS-bright spectra of more luminous IR galaxies at similar redshifts (C08). The fainter luminosity range studied here is particularly important, as it contains the bulk of the IR background energy at these redshifts.  For galaxies at $0.2<z<0.3$, we took advantage of the simultaneous presence of the \\ha and \\hb lines in the zCOSMOS spectra to study the relation between line emission, broad-band colours and dust extinction.   We also discussed the possibility for the IR emission to be associated with dust heated by old rather than young stellar populations.  We summarise our results as follows:  $\\bullet$   92\\% of our \\tfm galaxies at $0.2<z<0.3$ have \\ha EW$> 5 \\, \\rm \\AA$, which indicates  that the vast majority of IR sources at these redshifts are associated with significant on-going star formation.  In most of these sources, the relation between the gas obscuration, i.e. $L(\\rm H\\alpha)/L(\\rm H\\beta)$, and IR/UV attenuation can be reasonably explained with different reddening laws of common use in the literature. However, for up to 16\\% of our sample at $0.2<z<0.3$, the Balmer decrement is much larger than predicted by any of these laws. We believe that this is produced by a combination of an inhomogenous dust distribution and the fact that the zCOSMOS spectra only receive the light from a 1-arcsec width slit centred at the source.  The presence of a very obscured star-forming region towards the centre can produce very high $L(\\rm H\\alpha)/L(\\rm H\\beta)$ ratios.   We note that this effect should be much less important at higher redshifts, where the 1-arcsec width covered by the {\\em VIMOS} slit collects a more representative fraction of the total  galaxy light.  It is necessary to emphasize that, although the  broad-band SEDs of many low-redshift IR galaxies are dominated by old stellar populations, the optical spectra reveal the presence of on-going star formation in most cases. Thus, for these galaxies, one cannot conclude that the IR light is produced by  dust heated by old stars. More likely, the young stellar component is the main energy source for the IR emission.  Of course, this does not exclude the possibility that the old stars have a minor contribution to the dust heating.   $\\bullet$ Although there is no one-to-one relation between the \\ha and \\hb emission and the galaxy optical broad-band colours, around 80\\% of the objects with \\ha and \\hb EW $> 5 \\, \\rm \\AA$  are in the blue ($B-i<1.3$) cloud, and around 70\\% of those with \\hb EW $< 5 \\, \\rm \\AA$ are in the `green valley' or redder ($B-i \\geq 1.3$).   $\\bullet$  The location in the BPT diagram of our galaxies with \\ha EW $> 5 \\, \\rm \\AA$  suggests that $\\sim 22\\%$ of them are composite systems. However, in none of these galaxies can the presence of an active nucleus be confirmed with  the X-ray or  {\\em IRAC} data available in the COSMOS field. This fact suggests that, although an AGN component might be present in these galaxies, in any case is not dominant, and that aperture effects could also influence the fraction of sources identified as composites.   $\\bullet$  We identified only 5 galaxies in our sample at $0.2<z<0.3$ which have no spectral signs of star formation and, besides, have elliptical or lenticular morphologies. We believe that these are the only candidates  in our sample at $0.2<z<0.3$  for which most of the IR emission could genuinely be produced by dust heated by old rather than new stellar populations.   \\vspace{0.5cm}  The second part of our analysis has been devoted to galaxies at $z>0.5$. We used the zCOSMOS spectra to determine oxygen abundances, measure the  $\\rm D_n(4000)$ parameter and the average $\\rm H_\\delta$ EWs. Our results are the following:   $\\bullet$  79\\% of our galaxies with \\hb EW $> 5 \\, \\rm \\AA$ at $0.5<z<0.7$ have degenerate metallicity values. This is a consequence of estimating metallicities based on only three emission lines (through the $R_{23}$ parameter).  When we place these galaxies in a stellar mass-metallicity diagram, we see that the upper-branch solution is compatible with the mass-metallicity relation in the local Universe within the error bars. This implies that our data do not allow us to conclude on any evolution of the mass-metallicity relation between $z\\sim0$ and $z\\sim0.6$. Among the 21\\% of galaxies with non-degenerate metallicity measurements, we find that more than a half are well below the mass-metallicity relation. Over our total IR-galaxy population at $0.5<z<0.7$, they put a secure lower limit of 12\\% to the fraction of outliers to this relation. This percentage is quite smaller than that observed for LIRGs and ULIRGs by C08. However, the comparison of our results with those for optically-selected galaxies with similar stellar masses and well-determined metallicities suggests that the 12\\% fraction is indeed a lower limit.   $\\bullet$   In contrast to the results for more luminous IR galaxies, we find that our present sample at $0.6<z<1.0$ displays a wide range of $\\rm D_n(4000)$ values.  Still, $\\sim 60 \\, (80)$\\% of our sources have un-corrected (corrected) $\\rm D_n(4000)<1.2$.   $\\bullet$  A combined analysis of our present faint IR sample and previous C08 brighter sample at $0.6<z<1.0$ has allowed us to investigate the possible modes of star formation in IR galaxies. The burst-like mode of star formation is necessary to simultaneously explain the different average spectral properties of the most luminous IR galaxies ($L_{TIR} \\gsim 3\\times 10^{11} \\, \\rm L_\\odot$). Instead, for nearly 40\\% of the less luminous LIRGs and most luminous IR normal galaxies, a continuous, quiescent star-formation mode is possible. Realistically, the $L_{\\rm TIR}$ for the transition at $0.6<z<1.0$ should be considered to be $(3\\pm2)\\times 10^{11} \\, \\rm L_\\odot$, to take into account the errors in the determination of the total IR luminosities and the data binning.  Our findings are consistent with other recent results in the literature, which indicate that around a half of IR galaxies of intermediate luminosities at $z\\sim1$ are typical disks that do not show any sign of disturbance (see e.g. Elbaz et al.~2007). They could also be explained in the scenario proposed by theoretical studies that suggest that some galaxies could have had a star formation history characterised by steady gas flows since early times (Dekel et al.~2009).   In summary, our optical spectral analysis has clearly shown the diverse nature of the galaxies that make the bulk of the mid-IR background at $z<1$. And, even when star formation seems to be the most important driver of the IR emission at these redshifts,  IR galaxies are characterised by a complex variety of star formation histories."
    },
    "0910/0910.3539_arXiv.txt": {
        "abstract": "Young stellar systems orbiting in the potential of their birth cluster can accrete from the dense molecular interstellar medium during the period between the star's birth and the dispersal of the cluster's gas.  Over this time, which may span several Myr, the amount of material accreted can rival the amount in the initial protoplanetary disk; the potential importance of this `tail-end' accretion for planet formation was recently highlighted by \\citet{throop08}.  While accretion onto a point mass is successfully modeled by the classical Bondi-Hoyle-Lyttleton solutions, the more complicated case of accretion onto a star-disk system defies analytic solution.  In this paper we investigate via direct hydrodynamic simulations the accretion of dense interstellar material onto a star with an associated gaseous protoplanetary disk.  We discuss the changes to the structure of the accretion flow caused by the disk, and {\\em vice versa}.  We find that immersion in a dense accretion flow can redistribute disk material such that outer disk migrates inwards, increasing the inner disk surface density and reducing the outer radius.  The accretion flow also triggers the development of spiral density features, and changes to the disk inclination.  The mean accretion rate onto the star remains roughly the same with and without the presence of a disk.  We discuss the potential impact of this process on planet formation, including the possibility of triggered gravitational instability; inclination differences between the disk and the star; and the appearance of spiral structure in a gravitationally stable system. ",
        "introduction": "Generally, stars form in clusters \\citep[e.g.][]{lada03}, and our current understanding of star formation is beginning to incorporate the effects of the cluster environment on the process rather than treating star formation as a one-system event occurring in isolation.  A complete picture of planet formation may require a similar awareness of the surroundings; the possible influence of the cluster on planet formation is a logical consequence of clustered star formation.  Protoplanetary disks have lifetimes of $\\sim 1$--10 Myr without significant external UV irradiation \\citep[e.g.][]{currie07,currie09}.  In the Orion Nebula Cluster, 80\\% of stars have disks \\citep{smith05}, and even in environments like the Trapezium cluster over 10\\% of stars may have disks containing a minimum mass solar nebula (MMSN, $\\sim 0.01$ \\msun) of material \\citep{mann09}.  These disk lifetimes are comparable to or longer than the crossing time of a typical star-forming region, and thus the cluster environment may affect planet formation or young planetary systems.  Disruptive possibilities include the truncation of disks  by interactions with other cluster members \\citep{clarke93,heller95,boffin98} or dynamically altering an existing planetary system \\citep{laughlin98,zakamska04,malmberg09}\\footnote{Although the rates for these interactions are probably low compared to formation timescales \\citep{adams06}.}, and photo-evaporation which can remove disks on $10^5$--$10^7$ yr timescales \\citep{hollenbach94,johnstone98}.  More constructive effects may include triggering planetesimal formation by external UV illumination \\citep{throop05}, and adding material to the disk by accretion from the cluster's remnant ISM reservoir \\citep[][hereafter TB08]{throop08}  The effect that we explore in this work is the accretion of ISM material by a star--disk system.  Bondi--Hoyle--Lyttleton accretion by a star moving through a background medium may play a fundamental role in determining the final stellar mass \\citep{bonnell97, bonnell01a}, an idea that has seen a large amount of study over the last decade.  A less-explored question is the importance of accretion after a star has already assembled the majority of its mass, and the cluster is dominated by the gravitational dynamics of the stars rather than the potential and bulk velocities of the molecular gas.  TB08 performed simulations of this `tail-end' accretion, finding that as stars orbit through the depleting molecular gas in a cluster they can accrete around 1 MMSN per 1 Myr across a range of standard star-forming regions.  The observational constraints on disk lifetimes allow for this accretion to take place while a primordial protoplanetary disk is still present.  A plausible example of the interaction between a disk and the ISM is found around the star HD 61005 \\citep{hines07}, displaying what appears to be a wind-swept disk.  The TB08 study highlighted the potential importance of this process and suggested several possible measurable effects, including altering the structure of the protoplanetary disk, changing its orientation, or changing its composition.  It is these ideas that we explore here via direct hydrodynamic simulations. TB08 treated accretion by an analytic prescription applied post-simulation to {\\it n}-body experiments, and here we present the results of exploratory simulations involving a protoplanetary disk that encounters a dense ISM at supersonic speeds.  In TB08 the accretion is treated as occurring from a mean background density over $\\sim 10^6$ yr; the variation we consider is a star-disk system encountering a dense filament or clump, with a higher accretion rate over a shorter amount of time.  After describing our setup and presenting the results of our experiments, we discuss their implications for theories of planet formation. ",
        "conclusions": "As discussed above, we were constrained by computational necessity to consider a disk interacting with a dense molecular ISM, corresponding perhaps to a filament or core in a star-forming region.  The analysis of TB08 which motivated this work dealt with an ISM density reflecting the mean value of a region, much lower than this work ($10^3$--$10^6$ cm$^{-3}$), and over longer timescales (Myr versus kyr).  Because of the nature of the interactions, we believe that the main results of this work should carry over from the dense interactions we simulated to the lower density, longer timescale interactions from TB08.  If ram pressure had been the dominant driver of the changes to the disk morphology, this would not be the case, so interactions with a much denser region or at much higher velocities will be qualitatively different than these.  However, the changes were instead chiefly a result of the deposition of material with no azimuthal angular momentum onto the disk over timescales longer than (or comparable to) the orbital timescale.  These effects should remain broadly the same, whether the low specific angular momentum material is incorporated into the disk over 10 or 1000 orbital periods.  Likewise, if the angular momentum is deposited episodically (as in TB08) versus in one interaction as we have modeled, the final effect should mostly be a function of the total mass deposited. Thus the main result of this work, the redistribution of the outer disk material inward, should be robust over a wide range of ISM parameters.  A potentially significant complication that we did not consider is that the ISM in a real star-forming cloud will not be smooth, but instead filamentary and clumpy.  To some extent we have modeled this, as the star moved from a region of zero density into the dense region, and left after moving through it for $\\sim 3000$ yr (equivalent to a core or filament $\\sim0.01$ pc across).  However, the effect of substructure on smaller scales would be an interesting direction for future studies.  Density or velocity structure on scales comparable to the accretion radius could have an important effect on how the disk restructuring proceeds, introducing net angular momentum to the inflow and perhaps changing the disk structure more dramatically than the almost totally-smooth ISM we have considered.  Simulations of density and velocity gradients in pure BHL accretion \\citep{ruffert97,ruffert99} show that strong angular momentum gradients can result in small, short-lived disk-like structures as the flow accretes onto the star.  The interaction of such an accretion flow with an extant disk warrants further investigation.  Another issue that is unmodeled here is the potential effect of protostellar outflows on the accretion flow.  There are two main effects which could alter the flow; radiation pressure, and mechanical winds or jets.  The influence of stellar radiation on BHL accretion has been investigated by \\citet{edgar04a}, who showed that for stellar masses above around 10 \\msuns the accretion rate can be significantly diminished, but for lower mass stars like those in this paper radiative feedback should not present a large impediment to accretion.  The question of jets' and winds' effects on an accretion flow was briefly considered in TB08.  The ram pressure of a jet from a typical T Tauri star could be enough to reverse the BHL flow.  However, the opening angle of collimated outflows tends to be $\\sim20^\\circ$ at a distance of 20-50 AU, and smaller at larger radii \\citep{hartigan04}.  The solid angle directly affected by a jet is then less than about 1 sr at scales comparable to the accretion radius, and assuming that the flow intersecting the jet is reversed, this effect is on the order of 10\\%.  The zero-inclination case we simulated would be more greatly affected, as the jet would be directly opposing the accretion wake.  This would most directly alter the accretion rate onto the star, however, and the accretion onto the disk would be less impacted.  To our knowledge the interaction between a collimated outflow and a BHL accretion flow has not been modeled numerically, and this presents and interesting question for further work, but we estimate that it should be a secondary effect.  With these thoughts about the broader applicability of our models in mind, we discuss the main results of our work and their implications for disk evolution and planet formation:  {\\em Changes to the disk surface density profile and morphology}.  We have identified two effects that change the surface density profile of the disk: accretion onto the disk from the ISM flow, and the subsequent redistribution of the extant disk material due to absorbing low angular momentum material.  While related in that the accretion of ISM material leads directly to the redistribution of the existing disk, the latter effect is more important in terms of changes to the disk surface density profile.  With accretion alone the surface density profile would increase $\\sim$10--20\\% (see figure \\ref{surfacedensities}), while with the migration of outer material inwards the profile more than doubles across much of the disk relative to the initial value.  A possible consequence of this, pointed out in TB08, is the delivery of fresh material to regions of the disk that may have been depleted by planet formation, possibly ushering in a second era of gas accretion.  TB08 note that over several Myr accretion can deliver an amount of gas comparable to (or even greater than) the original disk mass.  This work suggests that the effect of accretion could be magnified by the presence of remnant gas disk, dumping not only the accreted ISM gas but the outer disk gas onto the inner disk and any planets that may be forming there.  The presence of a dissipative disk can have an important influence on the dynamics of multi-planet systems \\citep{moorhead08,moeckel08,thommes08,matsumura09}, and the dynamics of such systems in a disk that is undergoing induced global migration may prove interesting.  A secondary effect is the shrinking outer radius of the disk.  As long as the disk is accreting ISM material, the inward migration of the outer regions continues, as is seen in figure \\ref{surfacedensities}.  Any wispy outer portions of the disk are more susceptible to ram pressure, and thus a clean disk edge is maintained.  This effect would be most pronounced in disks that have undergone more accretion, perhaps due to being in a denser-on-average region, for instance the center of a cluster potential.  As the disk redistribution occurs on a timescale comparable to the orbital period, this process could take effect quickly and perhaps provides a mechanism to generate a relationship between disk-radius and the system's radius in a cluster.  This would work in the same sense as later photoevaporation once a cluster's O stars begin to dominate their locality, for example as in NGC2244 \\citep{balog07}.  \\begin{figure} \\centering \\plotone{FourierComps.eps} \\caption{Amplitudes of the first eight Fourier components in the azimuth of the disk surface density.  {\\it Top}: the face-on case, showing some power in modes with $m\\lesssim5$, dominated by $m=1$.  {\\it Bottom}: the $30^\\circ$ inclination case, with a stronger $m=1$ mode due to the misalignment between the inflow vector and disk symmetry axis.  Note the difference in scale between the two plots.} \\label{fourier} \\end{figure} {\\em Triggered gravitational instabilities?} With our chosen parameters, the disks should be very stable against gravitational instabilities that may aid planet formation \\citep[e.g.][]{boss97,boss03,rice06}.  Despite this, we see spiral features in the column density plots (figures \\ref{flowcomp0} and \\ref{flowcomp30}).  To better understand these, we examine the azimuthal Fourier amplitudes of the surface density, which we compute as \\begin{equation*} A_m = \\sum_{j=1}^{n_{disk}} m_{j}  e^{-i m \\phi_{j}} , \\end{equation*} where $m_j$ is the mass of the SPH particle, $m$ is the mode, $\\phi_{j}$ is the azimuthal angle of the particle, and the summation is over the particles identified by their angular momentum as part of the disk.  In figure \\ref{fourier} we plot the amplitude of the first eight Fourier modes (normalized by $A_0$) for each of the runs, at five times each.  The Fourier modes for the isolated disk, which we do not plot, lack any power above $|A_m|/|A_0| = 0.003$ in the first eight modes, confirming that absent the influence of the ISM flow, these disks are smooth and stable.  In the face-on case, the spiral structure begins with a marginal increase in the amplitude of the $m \\sim$ 4--5 modes which are tracing the spiral features most clearly seen in the top panel of figure \\ref{flowcomp0}.  The low amplitudes suggest that these spiral features are not globally organized features due to gravitational instability, but rather due to a more local reorganization of the disk as the outer material is redistributed inwards.  The more striking feature is the $m=1$ mode, which dominates the Fourier spectrum at later times for $i=0^\\circ$ and at all times for $i=30^\\circ$.  This is due to the star accreting a filament of material with some momentum perpendicular to the accretion axis, so that the star is knocked slightly out of the center of the disk.  As the disk responds to this off-axis potential, the $m=1$ mode is excited.  The signal is much stronger in the inclined case, as even accretion of material along the inflow direction will move the star slightly out of the disk's axis of symmetry.  Thus this signal, too, is not due to gravitational instability, and the spiral features are unlikely to be signals of disk fragmentation leading to planet formation.  We note, however, that a disk beginning in a more marginal state of stability than ours could be pushed into an unstable state by the disk redistribution we see.  If a factor of $\\sim$2 increase in surface density would be enough to push a disk into an unstable state, than interactions with the ISM like we have modeled here may be sufficient to stimulate planet formation by gravitational collapse.  For the low-mass disk simulated here though, we stress that even with the addition of mass from the ISM the Toomre Q parameter remains much larger than unity.  The spiral features are {\\em not} due to gravitational instability, but instead the interaction of the disk and the star with an accretion flow.  This process could conceivably masquerade as gravitational instability in observations.  {\\em Changes to the disk inclination}. Intuitively the addition of material with angular momentum misaligned with the rotation of the disk seems like a potential route toward the inclination difference between the ecliptic plane and the Sun's rotation axis.  The instantaneous orientation of the disk, measured by the average angular momentum of the particles identified as belonging to the disk as in the surface density calculation, is a bit noisy, but at the end of the simulations when the disk has settled down the face-on case exhibits a $0.5^\\circ$ inclination change relative to its initial orientation, and the inclined case has changed by $1.5^\\circ$.  This mechanism would appear to not be enough to account for the Sun's $\\sim7^\\circ$ inclination relative to the ecliptic plane \\citep{beck05}, at least with the amount of mass accreted in our simulations.  If the accreted ISM material overwhelms the initial disk, as shown to be possible in TB08, the situation might be different, but a much more detailed examination of the timescales for planet formation versus disk replacement and reorientation would be necessary to make this claim.  To conclude, we have performed simulations studying the interaction between a protoplanetary disk and  a dense, ambient ISM.  The accretion rate from the ISM onto the star is fairly well described by the standard BHL accretion rate, even in the presence of a disk altering the flow.  We have showed that these interactions can have a significant effect on the physical structure of the disk, with the main effect being the redistribution the outer disk inwards.  This redistribution, and the addition of foreign material to the extant disk, may have some implications for planet formation scenarios.\\\\"
    },
    "0910/0910.1016_arXiv.txt": {
        "abstract": "We are carrying out a spectroscopic monitoring of Galactic Wolf-Rayet stars, in order to detect binary systems. The sample consists of approximately 50 stars of the Nitrogen sequence (WN) and fainter than $V$=13. The observations are made from the 4-m telescope at CTIO, Chile. In the following, we present the first results of the 2007-2008 campaign. ",
        "introduction": "Las estrellas Wolf-Rayet (WR) de Poblaci\\'on I, tienen vientos estelares muy intensos, densos, y calientes, que le dan la forma caracter\\'istica a sus espectros, i.e. l\\'ineas de emisi\\'on intensas y anchas. Seg\\'un la presencia de distintas l\\'ineas en el espectro\\'optico, se clasifican en WN (Nitr\\'ogeno), WC (Carbono), y WO (Ox\\'igeno). Los par\\'ametros intr\\'insecos de las estrellas WR son muy dif\\'iciles de determinar dado que muchas de sus propiedades est\u00e1n enmascaradas por sus poderosos vientos.  Las estrellas WR son consideradas des\\-cen\\-dien\\-tes de las estrellas tipo O, aunque la forma en que se produce esta transici\\'on es a\\'un motivo de debate. Hay un subconjunto de las estrellas WN, designadas como WNH, que se las considera estrellas O muy masivas con fuertes vientos estelares en una etapa evolutiva a\\'un temprana, distintas de las WR ``cl\\'asicas'' que queman He en sus n\\'ucleos (Conti \\& Smith, 2008).  La masa estelar es un par\\'ametro astrof\\'isico fundamental cuyo conocimiento, junto al \\'indice de p\\'erdida de masa, la composici\\'on qu\\'imica, y la rotaci\\'on, permite determinar las propiedades y evoluci\\'on de las estrellas, en especial las m\\'as masivas (Meynet \\& Maeder, 2005). El m\\'etodo m\\'as directo para medir masas estelares es a trav\\'es de la aplicaci\\'on de la tercera ley de Kepler. Las masas t\\'ipicas determinadas en estrellas WR tienen valores entre 9--16\\,M$_\\odot$ para las WC y 10--83\\,M$_\\odot$ para las es\\-tre\\-llas WN (cf. Crowther 2007). Las candidatas a ser las estrellas m\\'as masivas de la V\\'ia L\\'actea parecen ser del tipo WNH, e.g. las 2 componentes WN6 de NGC3603-A1, con masas absolutas de 116 y 89\\,$M_\\odot$ (Schnurr et al. 2008), y la componente WN6 de WR~21a que result\\'o tener una masa m\\'inima de 87 M$_\\odot$ (Niemela et al. 2008). Figer (2005) determina por an\\'alisis estad\\'istico que se deber\\'ia esperar descubrir estrellas con hasta 150 M$_\\odot$, pero los objetos descubiertos al momento a\\'un no llegan a ese l\\'imite, por lo que se espera poder hallarlos, y confirmar (o no) entonces dicho valor.  La fracci\\'on de sistemas binarios detectados entre las estrellas WR de nuestra galaxia es de $\\sim40$\\% (van der Hucht, 2001), sin embargo hay al menos dos razones para sospechar que \\'este es un l\\'imite inferior. Una raz\\'on es dada por Langer \\& Heger (1999), quienes argumentan que la mayor\\'ia de las Supernovas tipo Ib/c ocurren en binarias WR+OB interactuantes, y dada la estad\\'istica derivada de las SN observadas se esperar\\'ia que existan m\\'as estrellas WR en sistemas binarios que simples. Otra raz\\'on es aportada por la observaci\\'on de estrellas del tipo WC. M\\'as del 80 \\% de las estrellas WC9 y WC8 parecen pertenecer a sistemas binarios (cf. Tuthill et al. 1999; Monnier et al. 1999). La pregunta obvia es si esta frecuencia alta de binarias es intr\\'inseca de este sub-tipo o puede ser aplicado a todas las estrellas WR. Esto, entonces, podr\\'ia implicar que la poblaci\\'on de estrellas WR est\\'a asociada a binaridad, a trav\\'es de la transferencia de masa durante una determinada fase de un sistema O+O, y que las estrellas WR son una clave en la evoluci\\'on de estrellas masivas.   Dadas estas consideraciones, estamos llevando a cabo un monitoreo espectrosc\\'opico de estrellas WR de nuestra galaxia, con el objeto de detectar variaciones de velocidad radial que indiquen la presencia de sistemas binarios. La muestra consiste de unas 50 estrellas de la secuencia del Nitr\\'ogeno (WN) y m\\'as d\\'ebiles que $V$=13. A continuaci\\'on, presentamos los primeros resultados de la campa\\~na 2007-2008. ",
        "conclusions": ""
    },
    "0910/0910.4891_arXiv.txt": {
        "abstract": "The concordance model of cosmology contains several unknown components such as dark matter and dark energy. Many proposals have been made to describe them by choosing an appropriate potential for a scalar field. We study four models in the realm of loop quantum cosmology: the Chaplygin gas, an inflationary and radiationlike potential, quintessence and an anti-Chaplygin gas. For the latter we show that all trajectories start and end with a type II singularity and, depending on the initial value, may go through a bounce. On the other hand the evolution under the influence of the first three scalar fields behaves classically at times far away from the big bang singularity and bounces as the energy density approaches the critical density. ",
        "introduction": "\\label{sec:intro} It is generally believed that our Universe started with an inflationary phase followed by a radiation and matter dominated era. However, classical cosmology is not able to tackle the problem of the initial conditions of the universe. One of the possible solutions to this problem is that our expanding Universe was preceeded by a contracting phase. But powerful singularity theorems based on classical general relativity (GR) forbid such a behavior unless one assumes a form of matter that violates the positive energy conditions or modified versions of gravity (see e.g. \\cite{Abdalla:05,Nojiri:08}). On the other hand it is believed that quantum gravity should solve this problem by generating ideal conditions for the genesis of our universe. Several proposals such as the pre-big bang string cosmology\\cite{Gasperini:03} and the ekpyrotic/cyclic models\\cite{Khoury:01,Steinhardt:02} modify dynamics with (perturbative) quantum gravitational effects but have so far had limited viability.  A generic nonsingular evolution through the big bang can only be achieved if nonperturbative effects of quantum gravity are incorporated. One of the leading nonperturbative background independent approach is loop quantum gravity (LQG)\\cite{Ashtekar:04,Rovelli:04,Thiemann:07}. One of the main predictions of LQG is that the underlying geometry is discrete. The application of the quantization methods of LQG to homogeneous spacetimes results in what is known as loop quantum cosmology (LQC) \\cite{Bojowald:08,Bojowald:00,Bojowald:00:2,Bojowald:00:3,Bojowald:02,Ashtekar:03,Bojowald:03}. The results of LQC not only provide new insights into the quantum structure of spacetime near the big bang singularity but also remove this singularity by extending the time evolution to negative times. It has been rigorously shown \\cite{Ashtekar:06,Ashtekar:06:2,Ashtekar:08} that the evolution of contracting semiclassical universes passes through the quantum regime and expands to semiclassical universes. This nonsingular bounce stems from the fact that the dynamics in LQC is governed by a discrete quantum difference equation in quantum geometry. On the other, it can be shown \\cite{Ashtekar:06,Ashtekar:06:2,Singh:05,Vandersloot:05} that an effective Hamiltonian on a continuum spacetime can be found which approximates well the quantum dynamics (for a Wheeler-DeWitt analog see \\cite{Nelson:08}). The modification arising from nonperturbative effects to the classical Friedmann equation includes a term $\\rho^2/\\rho_{\\mathrm{crit}}$, where $\\rho$ is the energy density and $\\rho_{\\mathrm{crit}}$ denotes the critical density of the order of magnitude of the Planck density. Since this term is negative the evolution bounces whenever the energy density reaches a density close to the Planck density.  The viability of the bounces for more general matter sources has been studied in e.g. \\cite{Singh:06,Nelson:07,Nelson:07:2, Cailleteau:08}, where it was shown that the behavior of solutions with inflationary and negative potentials are nonsingular, respectively, where solutions of exponential potentials are analyzed. Moreover it was shown that for negative potentials the inner boundary also appears, corresponding to the classical recollapse, which leads to solutions having cyclic behavior. In \\cite{Cailleteau:09,Chiou:07} the authors studied the role of LQC effects in the Ekpyrotic/Cyclic model in Bianchi type I models and showed that the universe undergoes multiple small bounces and the anisotropic shear remains bounded throughout the evolution.  In this work we are interested in potentials which play a major role in classical cosmology. We first introduce the effective dynamics in LQC in Sec.~\\ref{sec:effdyn}. In Sec.~\\ref{sec:Singularities} we give a short overview of conditions for singularities occuring in Friedmann-Robertson-Walker (FRW) cosmologies. In Sec.~\\ref{sec:chaplygin} we study the Chaplygin gas and in Sec.~\\ref{sec:inflrad} we study the robustness of the bounce for a scalar field which has the property of being inflationary at small times and radiationlike at later times \\cite{Bouhmadi:09:2,Bouhmadi:09,Bouhmadi:09:3}. Sec.~\\ref{sec:antich} is devoted to the anti-Chaplygin gas and Sec.~\\ref{sec:quintessence} to a quintessence model which models dark energy. We conclude with Sec.~\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion} As an attempt to solve the shortcomings of the concordance model of cosmology several models of scalar fields have been proposed which interpolate between two stages of the evolution of our universe. We studied the influence of three types of scalar fields of cosmological interest, nameley the Chaplygin gas, a modificated version of it and quintessence, and one more exotic type called the anti-Chaplygin gas. While the first type models a unification of dark matter and dark energy, the second one interpolates between an early inflationary phase and radiation. Contrary to quintessence which was introduced as an effort to describe dark energy in terms of a scalar field, the anti-Chaplygin can be considered as a toy model without any direct application to cosmology.  We presented the solutions to the Friedmann equations in LQC for these four models. We showed that the evolution of the first three models (Chaplygin gas, modified Chaplygin gas and quintessence) follows the classical path until it approaches the critical density, where the modification to the Friemann equation gains in importance. As this modification is negative definite the evolution bounces and the Hubble rate changes sign. Some time after the bounce the evolution follows once again the classical trajectory. We showed that, while all the origin in the $(\\phi,\\dot\\phi)$-phase diagram acts as an attractor for the Chaplygin gas, the solutions of the modified version converge toward $\\dot\\phi\\rightarrow0$ and $\\phi\\rightarrow\\pm\\infty$. Quintessence behaves in a similar way, except that there are two independent sectors in the $(\\phi,\\dot\\phi)$-phase diagram such that trajectories with positive respectively negative initial $\\phi$ always stay positive respectively negative. The situation is radiacally different for the anti-Chaplygin gas where every trajectory starts and ends with a Type II singularity. Depending on the initial data the evolution may go through a bounce, however LQC is, as expected, not able to remove these Type II singularities. Because of this very fact there are also two independent sectors for $\\phi$."
    },
    "0910/0910.1366_arXiv.txt": {
        "abstract": "Globular clusters (GCs) are spheroidal concentrations typically containing of the order of $10^5$ to $10^6$, predominantly old, stars. Historically, they have been considered as the closest counterparts of the idealized concept of ``simple stellar populations.'' However, some recent observations suggest than, at least in some GCs, some stars are present that have been formed with material processed by a previous generation of stars. In this sense, it has also been suggested that such material might be enriched in helium, and that blue horizontal branch stars in some GCs should accordingly be the natural progeny of such helium-enhanced stars. In this contribution we show that, at least in the case of M3 (NGC~5272), the suggested level of helium enrichment is not supported by the available, high-precision observations. ",
        "introduction": "Globular clusters (GCs) have long been thought to provide one of the closest approximations to the idealized concept of ``simple stellar populations,'' with all stars in a given GC having closely the same distance, age, and chemical composition. However, our understanding of these objects is quickly changing, thanks not only to measurements of chemical inhomogeneities, but also to improved color-magnitud diagrams (CMDs), which show that at least some GCs present multiple populations, presumably indicative of multiple formation episodes.  Besides the classical cases of $\\omega$~Cen (e.g., \\cite{Vetal07}) and M54 (e.g., Siegel et al. 2007), which host multiple red giant branches (RGBs), subgiant branches (SGBs), and main sequences (MSs), a triple MS has recently been identified in NGC~2808 (Piotto et al. 2007), whereas a double SGB has also been found in 47~Tuc (\\cite{Aetal09}), M22 (\\cite{Metal09}), NGC 1851 (\\cite{Metal08}), NGC 6388 (\\cite{P08}), and NGC 5286 (Piotto et al. 2009). Several of these CMD peculiarities often go hand-in-hand with multimodalities along the horizontal branch (HB).  As a rule, different formation episodes inside the clusters have been suggested as the explanation for this phenomenon. However, depending on the specific case considered, different chemical enrichment scenarios have been advocated, including a spread in the CNO elements (\\cite{Cass08}, \\cite{Pietri09}) or/and various populations with different initial He abundances. As to the latter, a helium abundance of 40\\% by mass has been found to be required in order to reproduce the bluest MSs in $\\omega$~Cen and NGC~2808 (\\cite{Petal05}, \\cite[2007]{Petal07}).  In addition to this, it has also been recently suggested that the spread in colors seen along the HB in GCs should be due to a spread in $Y$ within these objects, in which case helium abundance enhancements would be the rule, rather than the exception, among GCs (e.g., \\cite{CD07}, \\cite{DC08}). In this sense, these authors have found that the helium abundance among M3's blue HB stars should be enhanced, with respect to M3's red HB population, by at least $\\Delta Y = 0.02$. The main purpose of this contribution is to test this key prediction (see also \\cite{Cetal09} for more details).  \\begin{figure}[hu] \\begin{center} \\begin{minipage}[t]{0.5\\linewidth} \\raisebox{-1.5cm}{\\includegraphics[height=2.5in]{aarvalcarcefig1.eps}} \\end{minipage}\\hfill \\begin{minipage}[t]{0.475\\linewidth} \\caption{Comparison between the equivalent evolutionary sequences (EESs) for a fixed metallicity and two different helium abundances (black lines: $Y = 0.23$; gray lines: $Y = 0.28$).}\\label{fig1} \\end{minipage} \\end{center} \\end{figure}   \\firstsection ",
        "conclusions": ""
    },
    "0910/0910.2686_arXiv.txt": {
        "abstract": "Classical T Tauri stars are pre-main-sequence objects that undergo simultaneous accretion, wind outflow, and coronal X-ray emission. The impact of plasma on the stellar surface from magnetospheric accretion streams is likely to be a dominant source of energy and momentum in the upper atmospheres of these stars. This paper presents a set of models for the dynamics and heating of three distinct regions on T Tauri stars that are affected by accretion:  (1) the shocked plasmas directly beneath the magnetospheric accretion streams, (2) stellar winds that are accelerated along open magnetic flux tubes, and (3) closed magnetic loops that resemble the Sun's coronal active regions. For the loops, a self-consistent model of coronal heating was derived from numerical simulations of solar field-line tangling and turbulent dissipation. Individual models are constructed for the properties of 14 well-observed stars in the Taurus-Auriga star-forming region. Predictions for the wind mass loss rates are, on average, slightly lower than the observations, which suggests that disk winds or X-winds may also contribute to the measured outflows. For some of the stars, however, the modeled stellar winds do appear to contribute significantly to the measured mass fluxes. Predictions for X-ray luminosities from the shocks and loops are in general agreement with existing observations. The stars with the highest accretion rates tend to have X-ray luminosities dominated by the high-temperature (5--10 MK) loops. The X-ray luminosities for the stars having lower accretion rates are dominated by the cooler accretion shocks. ",
        "introduction": "Nearly all low-mass stars exhibit magnetic activity and some kind of mass outflow. Young stars add to this already complex situation by also undergoing {\\em accretion} along a subset of the magnetic field lines that connect the star and its surrounding disk. Out of the many unanswered questions regarding star formation and early stellar evolution, many of them share the need to better understand coronal heating and the three-dimensional magnetohydrodynamic (MHD) interplay between winds and accretion. These physical processes are key to determining how rapidly the stars rotate, how active the stars appear from radio to X-ray wavelengths, and how the stars affect nearby planets.  For classical T Tauri stars (CTTS), there are several possible explanations of how and where outflows arise, including extended disk winds, compact X-winds, impulsive eruptions, and stellar winds that bypass the accretion disk altogether (see, e.g., Paatz \\& Camenzind 1996; Ferreira et al.\\  2006; Shu et al.\\  2007; Casse et al.\\  2007; G\\'{o}mez de Castro \\& Verdugo 2007; Edwards 2009; Aarnio et al.\\  2009; Romanova et al.\\  2009; Krasnopolsky et al.\\  2009). It is of definite interest to determine how much of the observed mass loss can be explained solely with stellar winds, since outflows that are ``locked'' to the surface are capable of removing excess angular momentum from the star (Matt \\& Pudritz 2005, 2007, 2008). In addition, there appear be multiple source regions of ultraviolet (UV) and X-ray activity, such as flares, accretion shocks on the stellar surface, and closed loops that may resemble the Sun's coronal active regions (e.g., Feigelson \\& Montmerle 1999; Kastner et al.\\  2002; Stassun et al.\\  2006, 2007; Argiroffi et al.\\  2007; G\\\"{u}del et al.\\  2007b; Jardine 2008).  Much of the research devoted to improving our understanding of coronal heating and wind acceleration has been focused on the Sun. After many decades of high-resolution observations, analytic studies, and numerical simulations, it seems increasingly clear that closed magnetic loops in the solar corona are heated by intermittent reconnection events that are driven by the convective stressing of the loop footpoints (e.g., Aschwanden 2006; Klimchuk 2006). However, the open field lines that connect the Sun to interplanetary space are most likely energized by the dissipation of waves and turbulent motions that originate at the surface and then propagate outwards (Tu \\& Marsch 1995; Cranmer 2002; Suzuki 2006; Kohl et al.\\  2006). Self-consistent models of turbulence-driven coronal heating and solar wind acceleration have begun to succeed in reproducing a wide range of observations without the need for ad~hoc free parameters (Rappazzo et al.\\  2007, 2008; Cranmer et al.\\  2007; Cranmer 2009). Such progress on the solar front provides a fruitful opportunity to better understand the fundamental physics of coronal heating, accretion, and wind acceleration in young stars.  The goal of this paper is to continue the construction and testing of self-consistent models for the atmospheric heating and wind acceleration of CTTS. The starting point for this work is an existing methodology developed by Cranmer (2008) for computing the properties of T Tauri stellar winds. These models incorporated a new physical mechanism for how the mid-latitude accretion streams can influence the polar outflow. The inhomogeneous impact of accreting material was suggested to generate strong MHD turbulence and waves that propagate over the stellar surface to energize the launching points of the stellar wind streams. However, these models were created only for an idealized evolutionary sequence of a one solar mass ($M_{\\ast} = 1 \\, M_{\\odot}$) star, and they dealt only with the polar wind. Both the mass loss rates and the X-ray fluxes predicted by Cranmer (2008) were lower than typical observed values for T Tauri stars by at least an order of magnitude. It is clear that these models need to be adapted to a more ``real world'' situation that encompasses both (1) actual measured parameters of individual T Tauri stars, and (2) a more comprehensive treatment of both the open and closed magnetic field structures on the stellar surface.  Section 2 discusses the specific CTTS stars selected for this study and summarizes some of the most salient fundamental properties derived from observations. Section 3 outlines the methodology used to estimate the magnetic field geometries, the accretion stream dynamics, and the properties of impact-generated waves for each modeled stellar case. The next three sections describe how the plasma properties in three main accretion-driven surface regions are computed. Section 4 deals with the acceleration of the stellar wind from the polar regions. Section 5 deals with the localized ``hot spot'' heating at the accretion shocks. Section 6 deals with the distributed heating in low-latitude coronal loops. Then Section 7 compares the modeled X-ray luminosities from the accretion shocks and coronal loops to existing observations. Finally, Section 8 contains a brief summary of the major results and discusses how the process of testing and refining the theoretical models should be improved in future work. ",
        "conclusions": "The primary aim of this paper has been to explore and test a set of physical processes that may be responsible for producing stellar winds and coronal X-ray activity in accreting T Tauri stars. A key new aspect of this work is the self-consistent generation of plasma heating and wind acceleration (via MHD waves and turbulence) for a database of well-observed CTTS in the Taurus-Auriga region. Scheurwater \\& Kuijpers (1988) and Cranmer (2008) suggested that inhomogeneous accretion can give rise to enhanced MHD wave amplitudes across the stellar surface. In this paper, the consequences of these turbulent motions were studied further and applied to both wave-driven wind models and footpoint-shaking models for coronal loops. By and large, the predicted X-ray luminosities from a combination of accretion shocks and hot loops agrees reasonably well with existing observations. The resulting mass loss rates for polar CTTS winds tend to be lower than those derived from measurements. This indicates that disk winds or X-winds (i.e., outflows not rooted to the stellar surface) may be responsible for many of those measurements. However, in several of the modeled cases the stellar winds do appear to contribute significantly to the measured mass loss rates.  The models presented above have shown that it is plausible to assume that the accretion energy is sufficient to drive CTTS stellar winds and coronal X-ray emission. However, this cannot explain the even stronger X-ray emission from WTTS, which are believed to not be accreting at all. Even if the lower X-ray fluxes of CTTS were the result of absorption in the accretion columns (e.g., Gregory et al.\\  2007), the relatively unattenuated X-rays of WTTS should not be explainable by accretion-driven phenomena. Here it was assumed that the accretion-driven component of surface MHD turbulence dominates the intrinsic convection-driven component. Young, rapidly rotating stars are suspected of having more vigorous convective motions, higher internal dynamo activity, and more frequent magnetic flux emergence than older stars such as the Sun (e.g., Montalb\\'{a}n et al.\\  2004; Ballot et al. 2007; Holzwarth 2007; Brown et al.\\  2008; K\\\"{a}pyl\\\"{a} et al.\\  2009). Rapid rotation can also give rise to complex field-line motions that lead to propeller-driven outflows, collimated jets, or ``magnetic towers'' (Blandford \\& Payne 1982; Lynden-Bell 2003; Ustyugova et al.\\  2006; Romanova et al.\\  2009). More accurate models must include these kinds of effects.  In this paper, an attempt was made to reduce the number of freely adjustable parameters by as much as possible. For example, although the accretion shock model (Section 5) depends on many assumptions about the magnetic geometry and the form of the radiative cooling function, it does not contain any truly ad~hoc free parameters. The coronal loop model (Section 6) utilizes a state-of-the-art parameterization for the heating that comes from MHD turbulent dissipation, and the best value for the $\\alpha$ exponent in that parameterization turns out to be one that has been fit from a series of numerical simulations (e.g., equation [\\ref{eq:alrapp}]). However, all of the models presented in this paper do depend on various assumptions that had to be made about the clumpiness of the accreting gas, and about how these clumps may be decelerated as they fall from the inner edge of the disk. Future models of inhomogeneous accretion should focus on improving our understanding of the dynamics of individual clumps.  The use of a rotationally aligned dipole magnetic field in this paper is a shortcoming that should be remedied in future models. In fact, there has been much work done recently to create three-dimensional models of CTTS magnetospheres (e.g., Koldoba et al.\\  2002; Romanova et al.\\  2003, 2004, 2008; Gregory et al.\\  2006, 2009; Long et al.\\  2007, 2008). Such models can naturally include the effects of X-ray absorption by the accretion columns and the dynamics of disk winds or X-winds. Still, many of these models utilize relatively simple approximations for the conservation of internal energy. Their realism and predictive power could be improved by including the effects of MHD waves and turbulence as described above. Thus, it is possible that the most useful results of this paper may not be the specific model results, but instead the general {\\em methodology} for wave generation and coronal heating that can be inserted into more advanced simulations.  In addition to improving the physical models of the stellar plasmas, it will also be important to model the emergent radiation more accurately. For X-rays, the simulation of bandpass-integrated luminosities (Section 7) is only the beginning of the story. There is a great deal of diagnostic information in the X-ray spectrum that can be used to constrain the properties of the various plasma components. Recent X-ray spectral analyses are beginning to reveal that there may be {\\em other} distinct emitting regions in addition to the three studied in this paper (e.g., winds, accretion shocks, and coronal loops). One of these other classes includes transient phenomena such as strong flares, eruptions, and other outcomes of large-scale magnetic reconnection (e.g., Giardino et al.\\  2006; Kastner et al.\\  2006; Aarnio et al.\\  2009). Another newly discovered region appears to be filled with cooler and lower-density coronal plasma than is in the dominant coronal loops (Brickhouse et al.\\  2009). These regions may be turbulent ``boundary layers'' that surround the accretion spots and are magnetically connected to both the shocked plasma and the loops (see also Hujeirat \\& Papaloizou 1998).  \\hspace*{0.01in}  I gratefully acknowledge Nancy Brickhouse, Adriaan van Ballegooijen, and Andrea Dupree for many valuable discussions, as well as Sean Matt for helping me discover an error in an earlier version of this paper. This work was supported by the Sprague Fund of the Smithsonian Institution Research Endowment, and by the National Aeronautics and Space Administration (NASA) under grant {NNG\\-04\\-GE77G} to the Smithsonian Astrophysical Observatory."
    },
    "0910/0910.2365_arXiv.txt": {
        "abstract": "We obtain analytical expressions for the velocity anomaly due to the Rossiter-McLaughlin effect, for the case when the anomalous radial velocity is obtained by cross-correlation with a stellar template spectrum.  In the limit of vanishing width of the stellar absorption lines, our result reduces to the formula derived by \\citet{Ohta2005}, which is based on the first moment of distorted stellar lines. Our new formula contains a term dependent on the stellar linewidth, which becomes important when rotational line broadening is appreciable.  We generate mock transit spectra for four existing exoplanetary systems (HD17156, TrES-2, TrES-4, and HD209458) following the procedure of \\citet{Winn2005}, and find that the new formula is in better agreement with the velocity anomaly extracted from the mock data. Thus, our result provides a more reliable analytical description of the velocity anomaly due to the Rossiter-McLaughlin effect, and explains the previously observed dependence of the velocity anomaly on the stellar rotation velocity. ",
        "introduction": "}  Among approximately 400 extrasolar planets discovered as of September 2009, there are more than 60 transiting planets. The transiting planets provide important information, such as radii and atmospheric signatures, which are unavailable from radial-velocity data alone. Another advantage of transiting exoplanets is that one can measure the angle between the stellar rotation axis and the orbital axis of the exoplanet projected onto the sky (conventionally denoted by $\\la$) through the Rossiter-McLaughlin (RM) effect. The RM effect generates a radial velocity anomaly during a transit, due to the asymmetric line profiles that result from the partial occultation of the rotating stellar disk \\citep{Rossiter1924, McLaughlin1924}. By carefully investigating velocity anomalies during transits, one can determine the trajectory of planets on the stellar disk and estimate $\\la$ precisely \\citep[e.g.,][]{Queloz2000,Ohta2005,Winn2005}.  This angle is especially important in understanding the basic physical processes of planetary formation and subsequent orbital migration. According to recent planetary formation theories, gaseous planets orbiting within 0.1 AU of the parent star (hot-Jupiters) are supposed to have formed a few AU away from the star, and subsequently to have migrated inward.  While standard migration mechanisms keep $\\la\\sim 0^\\circ$, scenarios such as planet-planet scattering may change $\\la$ significantly \\citep{Lin1996,Chatterjee2008,Nagasawa2008}.  The Kozai mechanism may also lead to a highly inclined orbit \\citep{Wu2007}. Thus, the observed distribution of the angles provides an observational clue to distinguish between different planet formation and migration theories.  \\citet{Ohta2005} derived an analytic formula (hereafter, the OTS formula) for velocity anomalies due to the RM effect, by computing the first moment (intensity-weighted mean wavelength) of distorted absorption lines perturbatively.  (Previously, \\citet{Kopal1942} and others used the first-moment method for eclipsing binary stars to estimate the velocity anomalies.)  The OTS formula proved to be useful in understanding the parameter dependence of the velocity anomaly and also in forecasting the error budget of the parameter estimate, in particular of the spin-orbit misalignment angle $\\lambda$. Indeed their work inspired \\citet{Winn2005} to revisit the estimate of $\\lambda$ precisely for the first discovered transiting planetary system HD~209458.  For that purpose, \\citet{Winn2005} generated in-transit stellar spectra of HD~209458, put them into the Keck analysis routine of radial velocities, and found that the OTS formula systematically underpredicts the velocity anomaly by 10\\%. Intriguingly, however, \\citet{Johnson2008} and \\citet{Winn2008}, found the OTS formula to provide an adequate description of the simulated results for the cases of HAT-P-1 and TrES-2.  We noticed that the systematic difference between the OTS formula and the simulation may be sensitive to the stellar spin rotation velocity $v\\sin i_\\star$ ($i_\\star$ denotes the inclination angle of the stellar spin axis), which is $4.70 \\pm 0.16~ {\\rm km} ~{\\rm s}^{-1}$, $3.75 \\pm 0.58~ {\\rm km}~ {\\rm s}^{-1}$, and $1.0 \\pm 0.6~ {\\rm km} ~{\\rm s}^{-1}$, for the HD~209458, HAT-P-1, and TrES-2 systems, respectively. In addition, the actual radial velocity measurement algorithm does not adopt the moment method, strictly speaking.  For instance, the analysis routines used by the HARPS (the High Accuracy Radial velocity Planet Searcher) and SOPHIE (Spectrographe pour l'Observation des Ph$\\mathrm{\\acute{e}}$nom$\\mathrm{\\grave{e}}$nes des Int$\\mathrm{\\acute{e}}$rieurs stellaires et des Exoplan$\\mathrm{\\grave{e}}$tes) teams involve the cross-correlation of the observed spectrum with a template spectrum, and the velocity anomaly is estimated from the peak of the cross-correlation function. The analysis routine which uses the iodine cell technique, with the Subaru HDS (the High Dispersion Spectrograph) or Keck HIRES (the High Resolution Echelle Spectrometer), involve a fit to the observed spectra, based on a template spectrum and the known spectrum of the iodine cell.   For these reasons, in this paper we revisit the analytic approach to the RM effect using the cross-correlation method.  Our analytic approach to this problem is complementary to the recent numerical approach by \\citet{Triaud2009}, and provides an analytic framework for understanding their results.  Even though the cross-correlation method is not directly applicable to the analysis of the iodine cell technique, we find that our result seems to capture the main qualitative features of the numerical results based on that technique; specifically, our formula includes a term dependent on the stellar spin velocity that does not show up in the moment method but was empirically found by \\citet{Winn2005}.   The rest of the paper is organized as follows. Section 2 describes our analytical modeling of stellar absorption lines distorted by the occultation due to a transiting planet (\\S 2.1) and derives the RM velocity anomaly on the basis of the moment method (\\S 2.2) and the cross-correlation method (\\S 2.3).  In order to check the reliability of the new analytic formula, we generate simulated spectra during transits and put them into the Subaru analysis routine of radial velocities in \\S 3 as \\citet{Winn2005} did for the Keck routine. Specifically we consider HD~17156 ($v\\sin i_\\star =4.2~ {\\rm km}~ {\\rm sec}^{-1}$), TrES-2 ($v\\sin i_\\star =1.0 ~{\\rm km}~ {\\rm sec}^{-1}$), TrES-4 ($v\\sin i_\\star =8.5 ~{\\rm km}~ {\\rm sec}^{-1}$), and HD209458 ($v\\sin i_\\star =4.5 ~{\\rm km}~ {\\rm sec}^{-1}$) systems, and find that our new formula reproduces the simulated data better than the OTS formula.  Finally, section 4 is devoted to a summary and further discussion of the present paper. ",
        "conclusions": "}  We have presented an improved analytic formula for the RM effect as measured from the cross-correlation method. Our main finding is equation (\\ref{gaussianDv}) which describes the RM velocity anomaly under the Gaussian approximation for the intrinsic line profile and the stellar rotation kernel. Unlike the previous approximation (the moment method) adopted by OTS, the radial velocity anomaly from the cross-correlation method is explicitly dependent both on the thermal and natural broadening width of an individual line profile ($\\beta$) and on the rotation velocity ($\\sigma$).  The analytical formula has been compared with the simulated analysis for the Subaru HDS routine using the mock spectra constructed from the high-resolution spectrum of the Sun. Even though the actual analysis routine attempts to fit simultaneously many different lines that are not necessarily approximated by Gaussians, the resulting RM velocity anomaly is well described by a form of equation (\\ref{gaussianDv}).  The current result explains the previous findings \\citep{Winn2005,Johnson2008,Winn2008} that the OTS formula does not provide a good approximation to planetary systems with a relatively high stellar rotation rate. The next step in obtaining a more accurate analytic formula for the RM effect would probably be to use a more accurate rotation kernel (eq.~[\\ref{rotation}]) instead of the Gaussian approximation (\\ref{eq:RapproxG}). For future work we plan to pursue this approach, together with the propagation of the resulting precision on the misalignment angle $\\lambda$ and the rotation velocity $v\\sin i_\\star$.  In addition to the improvement of the analytic approach, the current result points to a possible further refinement of the observational reduction routine for the velocity anomaly due to the RM effect. The procedure described in \\S 3.1 is appropriate for the radial velocity determination {\\it out of transit}. During the transit, however, the Doppler-shifted stellar spectrum $S(\\la-\\d\\la)$ in equation (\\ref{model}) should be replaced with \\begin{eqnarray} &S_{\\mathrm{star}}(\\la-\\tilde{\\la})-S_{\\mathrm{planet}}(\\la-\\tilde{\\la}-\\D\\la) \\nonumber\\\\ =&S_{\\mathrm{star}}(\\la-\\tilde{\\la}) -fS_{\\mathrm{narrowed~ template}}(\\la-\\tilde{\\la}-\\D\\la), \\end{eqnarray} where $S_{\\mathrm{star}}(\\la)$ is the stellar template spectrum derived by the procedure described in \\S3.1 during the out-of-transit data, $S_{\\mathrm{narrowed~ template}}(\\la)$ is obtained by deconvolving $S_{\\mathrm{star}}(\\la)$ with the rotational kernel in principle, and $\\tilde{\\la}$ and $\\D\\la$ correspond to the Keplerian plus peculiar velocity of the star and the radial velocity of the occulted portion of the stellar disk, respectively.  Fitting the observed spectrum {\\it in transit} with the additional parameters $f$, $\\tilde{\\la}$, and $\\D\\la$, instead of $\\d\\la$ alone, one may directly extract the velocity $v_p= c\\D\\la/\\la_0$ after averaged over many different lines.  Moreover one may reduce the number of fitting parameters by estimating $f$ from photometric data during the transit.  Since the above method is more appropriate to model the spectra of the transiting planetary system during transit, one expects that the precision and accuracy of $\\D\\la$, and therefore the RM velocity anomaly, may be improved in principle.  In reality, however, the practical feasibility to implement in the analysis routine depends on the available resolution of the spectra as well as their stability. The work toward this improvement is also currently under way, which we also hope to present elsewhere in due course.  The RM effect has become an observationally mature probe of the transiting planetary systems. Further improvement in both the accuracy and the precision of the velocity anomaly is important to advance the understanding of the formation and evolution processes of such systems. We hope that the present analytic approach provides some insights that will enable progress in that direction."
    },
    "0910/0910.2709_arXiv.txt": {
        "abstract": "We present the first results from the \\sw\\ survey, an ongoing project to identify compact white dwarf (WD) binaries in the spectroscopic catalog of the Sloan Digital Sky Survey. The first object identified by \\sw, \\WDT, is a single-lined spectroscopic binary in a circular orbit with a period of $4.56$ hr and a semiamplitude of $322.7\\pm6.3\\,\\mathrm{km\\,s^{-1}}$. From the spectrum and photometry, we estimate a WD mass of $0.92^{+0.28}_{-0.32}\\,\\mathrm{M_{\\odot}}$. Together with the orbital parameters of the binary, this implies that the unseen companion must be more massive than $1.62^{+0.20}_{-0.25}\\,\\mathrm{M_{\\odot}}$, and is in all likelihood either a neutron star or a black hole. At an estimated distance of $48^{+10}_{-19}$ pc, this would be the closest known stellar remnant of a supernova explosion. ",
        "introduction": "\\label{sec:Intro}  The vast majority of stars with masses below $8 \\mathrm{M_{\\odot}}$ end their lives as white dwarfs (WDs). Left to their own devices, WDs will just cool and fade away, but WDs that form part of close binary systems can interact with their companions, steering away from the normal course of isolated stellar evolution and giving rise to a plethora of astrophysical phenomena. When the companion of the WD is another compact object (a second WD, a neutron star [NS], or a black hole [BH]), we refer to these systems as compact WD binaries, or CWDBs. The orbital evolution of detached CWDBs is governed by the emission of gravitational waves \\citep[GWs,][]{peters63:GW_radiation_Keplerian_orbit,paczynski67:GW_close_binaries}. For systems with small enough separations, loss of angular momentum leads to a merger of the two components in a time that is short enough to be of astrophysical interest. This kind of CWDBs constitute the largest population of Galactic GW sources \\citep{evans87:WD_binaries_as_GW_sources}, dominating the GW foreground for future missions like LISA \\citep{nelemans09:binaries_review}. They can also provide important constraints on the still poorly understood process of common-envelope evolution in binary systems \\citep{nelemans05:common_envelope_in_WD_binaries}.  But perhaps the most interesting aspect of CWDBs is the outcome of the merging process at the end of their binary evolution. For systems composed of two CO WDs whose combined mass is above the Chandrasekhar limit, the merger might trigger a thermonuclear runaway and lead to a Type Ia supernova (SN) explosion \\citep{iben84:typeIsn,webbink84:DDWD_Ia_progenitors}. This possibility, often referred to as the double degenerate WD (DDWD) SN Ia progenitor scenario, is the only theoretical model that naturally explains the absence of H in the nebular spectra of Type Ia SN \\citep[see][and references therein]{Leonard07:H_nebular_Ia_spectra}, and it has motivated a number of searches for suitable candidate systems. The most comprehensive of these searches, the SPY survey \\citep{napiwotzki01:SPY_survey}, examined the spectra of $\\sim1000$ WDs to look for radial velocity (RV) shifts in the characteristic absorption lines. The total number of detached DDWDs found to date by SPY and other surveys is around 100 \\citep{napiwotzki04:SPY_04}, but the periods and mass estimates have only been published for 24 systems \\citep{nelemans05:SPY_IV}. So far, none of these systems clearly fulfills the requisites to be a SN Ia progenitor.  In this paper and in a companion publication \\citep{mullally09:DDWDs}, we present the first results from \\sw, the Sloan White dwArf Radial velocity data Mining Survey. The aim of \\sw\\ is to mine the WD catalog in the spectroscopic data base of the Sloan Digital Sky Survey \\citep[SDSS,][]{york00:SDSS_Technical} in search of CWDBs. Our ultimate goal is to find the DDWD progenitors of Type Ia SNe, or at least put constraints on the rate of DD mergers in the Galactic disk that can be compared with measurements of the local Type Ia SN rate. This will allow us to assess the viability of the DDWD progenitor scenario for SN Ia. This paper is organized as follows. Our data mining strategy is briefly described and compared to the SPY survey in Section \\ref{sec:Strategy}, using as an example the CWDB \\WDTW\\ (henceforth, \\WDT), the first object discovered by \\sw. In Section \\ref{sec:Followup}, we present the follow-up observations of \\WDT\\ that we have conducted in order to determine its orbital and spectral parameters. In Section \\ref{sec:Companion}, we discuss the implications of our findings for the nature of the unseen companion of \\WDT. We summarize our conclusions and outline the future prospects for our survey in Section \\ref{sec:Conclusions}. ",
        "conclusions": "\\label{sec:Conclusions}  In this paper and a companion publication \\citep{mullally09:DDWDs} we have presented the first results from the SWARMS survey, an ongoing project aimed at discovering and characterizing CWDBs in general and DDWD SN Ia progenitors in particular among the WDs in the spectroscopic SDSS data base. Further results and a complete description of the survey will be the subject of forthcoming publications. Our ultimate goal is to estimate a rate for the merger of DDWD SN Ia progenitors in the Galaxy, and to use that rate, in combination with the results from SPY and other surveys, to assess the viability of the DDWD progenitor scenario for Type Ia SNe.  To demonstrate the capabilities of SWARMS, we have focused on \\WDT, the first object found by the survey. This WD was identified as a member of a binary system from the SDSS spectra, and follow-up observations were performed to measure its orbital and spectral parameters. The RV curve has a period of $4.5550 \\pm 0.0007$ hr, with a semiamplitude of $322.7 \\pm 6.3\\,\\mathrm{km\\,s^{-1}}$. The primary WD is cold ($\\sim9000$ K), which makes the spectral analysis somewhat challenging, but we have derived a conservative estimate of $0.92^{+0.28}_{-0.32}\\,\\mathrm{M_{\\odot}}$ for its mass. This implies that the unseen companion must be larger than $1.62^{+0.20}_{-0.25}\\,\\mathrm{M_{\\odot}}$, and is probably a NS or a BH. At a distance of $D=48^{+10}_{-19}$ pc, this would be the closest remnant of a SN explosion to the Solar System. We have discussed the implications of our discovery, and suggested future avenues of research on this object. This kind of massive, close binary is a perfect example of the objects that can be found by SWARMS and other similar surveys like MUCHFUSS \\citep{tillich09:MUCHFUSS}. We anticipate more such discoveries, which we will report in the literature. We also plan to set up an on-line database of SWARMS objects."
    },
    "0910/0910.5635_arXiv.txt": {
        "abstract": "It is shown that the standard cosmological parallax--distance formula, as found in the literature, including text-books on cosmology, requires a correction. This correction arises from the fact that any chosen baseline in a gravitationally bound system does not partake in the cosmological expansion and therefore two ends of the baseline used by the observer for parallax measurements cannot form a set of co-moving co-ordinates, contrary to what seems to have been implicitly assumed in the standard text-book derivation of the parallax distance formula. At large redshifts, the correction in parallax distance could be as large as a factor of three or more, in the currently favoured cosmologies (viz. $\\Omega_\\Lambda=0.73, k=0$). Even otherwise, irrespective of the amount of corrections involved, it is necessary to have formulae bereft of any shortcomings. We further show that the parallax distance does not increase indefinitely with redshift and that even the farthest observable point (i.e., at redshift approaching infinity) will have a finite parallax value, a factor that needs to be carefully taken into account when using distant objects as the background field against which the parallax of a foreground object is to be measured. ",
        "introduction": "Cosmological distance formulae are standard formulae for different world models in use for more than half a century and are available in review articles and text books on cosmology. The two most commonly used formulae for cosmological distant sources are the luminosity distance formula and the angular diameter distance formula. A use of either of these in conjunction with the measured values of flux densities or angular sizes of a suitable sample, for ascertaining the geometry of the universe, depends upon the assumption of the existence of a standard candle/rod, and a successful interpretation is very much dependent on the cosmological evolution of the intrinsic luminosities/physical sizes of the parent population of the sources of interest. Not only the large spread in their intrinsic values makes it very uncertain to define a standard candle/rod, in fact in most cases the effects of the evolutionary changes in source properties are overwhelmingly larger than those expected due to differences in geometry between different world models. The other two formulae, which presently are being used only for nearby objects, are the proper motion distance formula and the parallax distance formula.  For testing different world models, application of the parallax distance in particular, unlike other distance measures, does not depend upon any assumption about the intrinsic properties of the observed sources. Thus it is also independent of the chosen frequency band since no source property is involved, as long as the source is detectable in that band. At a given redshift, the observed parallax depends only on the chosen baseline of the observer and the world--model geometry.  Weinberg (1970) showed that even otherwise, the measurement of redshift and luminosity (or angular diameter) distance cannot in principle determine the sign and magnitude of the spatial curvature unless supplemented with a dynamical model. However, this ambiguity can be resolved by parallax measurements at cosmological distances. With the achieved angular resolutions already in the micro-arcsec domain (Fomalont \\& Kobayashi 2006), it could with the advancement of technology become important rather sooner than expected.  The cosmological parallax--distance formulation is available in text-books on cosmology (Weinberg 1972; Peacock 1999), review articles (von Hoerner 1974) or reference books (Lang 1980), with expressions for the different world models given there. Here we show that a subtle correction is required in the standard formulation found in the contemporary literature. The correction stems from the fact that the physical dimensions of a gravitationally bound system like that of the solar system (or for that matter even of larger systems like that of a galaxy), do not change with the cosmological expansion of the universe. Therefore two ends of a baseline, used by the observer for parallax measurements, do not partake in the free-fall-like motion of the cosmic fluid and thus cannot be considered to form a set of co-moving co-ordinates, contrary to what seems to have been implicitly assumed in the standard text-book derivations of the parallax distance formula. The correction can be very large (a factor of three or even more) at large redshifts. In any case, irrespective of the amount of corrections involved, it is imperative to have formulae bereft of any shortcomings. In the next section we shall briefly review the formulation given in the literature and in the section following that, we spell out the required corrections. ",
        "conclusions": ""
    },
    "0910/0910.1095.txt": {
        "abstract": "We report multi-wavelength power spectra of diffuse Galactic dust emission from BLAST observations at 250, 350, and 500~\\micron\\ in Galactic Plane fields in Cygnus~X and Aquila. These submillimeter power spectra statistically quantify the self-similar structure observable over a broad range of scales and can be used to assess the cirrus noise which limits the detection of faint point sources.  The advent of submillimeter surveys with the \\textit{Herschel Space Observatory} makes the wavelength dependence a matter of interest. We show that the observed relative amplitudes of the power spectra can be related through a spectral energy distribution (SED).  Fitting a simple modified black body to this SED, we find the dust temperature in Cygnus~X to be \\cpt\\ K and in the Aquila region \\apt\\ K. Our empirical estimates provide important new insight into the substantial cirrus noise that will be encountered in forthcoming observations. ",
        "introduction": "\\label{sec:intro} Energy is being injected continually into the interstellar medium (ISM) through spiral shocks, violent outflows from massive protostars, stellar winds, expanding \\ion{H}{2} regions and supernova explosions. The result is a turbulent medium where the dust is well mixed in the structures produced.  The emission from the relatively nearby ISM at high Galactic latitude is known as Galactic Cirrus.  Cirrus-like structure in the brighter emission near the Galactic Plane has been called `interstellar froth' \\citep{waller1994}; this might be an interesting distinction with a physical basis, but here we will simply use the term `cirrus'.  The dynamics of the ISM appear to have made the distribution of density structures self-similar, with fluctuations in column density and surface brightness present on all observable scales, though decreasing towards smaller scales.  Statistical description of random fluctuations through structure functions is an important method of extracting physical properties hidden in diffuse emission as well as for providing a quantitative measure to compare with simulations. % A common statistical tool used to estimate the level of cirrus noise is the power spectrum.  This is valid for Gaussian random fields. However, \\citet{gautier} found evidence for non-Gaussianity and recent studies by \\citet{ma-dust} have revealed non-vanishing skewness and excess kurtosis in the underlying brightness fluctuation fields. Nevertheless, for estimating the variance the power spectrum is still indicative.  Using power spectra, \\citet{gautier} quantified how fluctuations associated with Galactic cirrus are a source of confusion noise which limits the detectability of point sources. Even with the improved angular resolution of imagers like SPIRE and PACS on the \\textit{Herschel Space Observatory}\\footnote{http://www.esa.int/SPECIALS/Herschel/index.html} (hereafter \\textit{Herschel}), this `cirrus noise' remains important, often dominant. Therefore, to assess the noise it is vital to know the statistical properties of the interstellar diffuse emission at the relevant wavelengths.  A number of other statistical analyses of emission have been carried out on the basis of power spectra for high latitude clouds, e.g., by \\citet{kiss-galcir}, \\citet{jeong2005}, and \\citet{ma-dust}. % In this paper we analyse diffuse dust emission in the Galactic Plane and present for the first time multi-wavelength power spectra in the submillimeter, based on observations in Cygnus~X (Cyg~X) and Aquila with BLAST \\citep{pascale2008} at 250, 350, and 500~\\micron. % We also analyse the same regions at 100 and 60~\\micron\\ using IRIS, \\IRAS\\ data reprocessed by \\citet{mairis}. % Compared to high latitude studies, analysing diffuse emission in the Galactic Plane in terms of the structure of the ISM resulting from its turbulent properties is more challenging because of the long path lengths, high column density, high star-formation rate, and contamination by compact sources. In addition, in the Galactic Plane, self-gravity can play an important role in shaping diffuse structures at smaller scales. % Of particular importance are the potentially strong spatial variations of the radiation field due to star formation activity.  As discussed further in \\S~\\ref{subsec:masscol}, this could produce variations in the dust emission independently of any changes in structure of the ISM.  Our paper is organized as follows. % % We begin with a brief description of BLAST observations (\\S~\\ref{sec:obs}) and then introduce important aspects of the power spectrum and accompanying cirrus noise (\\S~\\ref{sec:ps}). % We analyse IRIS data in \\S~\\ref{sec:resI}, placing this in the context of earlier studies and providing a short-wavelength reference for our submillimeter studies. % In \\S~\\ref{sec:resB} we examine the BLAST data: the noise; effect of the point spread function (PSF) or beam; removal of compact sources; and the exponents and amplitudes of the submillimeter power spectra. We estimate the cirrus noise for these maps and compare it with the completeness depth of the BLAST05 point-source catalogs at 250~\\micron. In \\S~\\ref{sec:depth} we also discuss some implications for related approved observations with \\textit{Herschel}. % We show in \\S~\\ref{sec:dis} how the observed wavelength dependence of the amplitude of the power spectrum can be understood as a straightforward consequence of the spectral energy distribution (SED) of the dust and we fit a simple modified black body to estimate the dust temperature. % Our empirical results provide new insight into what cirrus noise might be expected in submillimeter observations. ",
        "conclusions": "\\label{sec:con}  In these fields in the Galactic Plane, the exponent of the power spectrum of the 100~\\micron\\ IRIS map is close to $-3$, within the dispersion seen by \\citet{ma-dust} despite the average brightness $\\langle I_{100} \\rangle$ being beyond the range studied by these authors. % On the other hand, the amplitudes of the 100~\\micron\\ power spectra estimated for these bright $\\sim 2$\\degree $\\times$ $2$\\degree\\ fields are significantly below what would be extrapolated from the trend with $\\langle I_{100} \\rangle$ found by \\citet{ma-dust}.  Therefore, particularly in such bright star forming regions, it is recommended that the power spectrum be computed directly. % The power spectra derived from the BLAST observations are also well fit by power laws, with similar exponents. % The frequency dependence of the amplitude of the power spectrum can be described by the square of an SED which is a simple modified black body function with a reasonable characteristic temperature. % This is confirmed by direct correlations between the maps at different wavelengths. % However, this characteristic temperature does appear to change in different Galactic environments, and unless its value is known independently the power spectra and/or map correlations need to be evaluated for all wavelengths.  Cirrus noise will be important in many planned multi-wavelength Galactic Plane, high latitude, and extragalactic surveys carried out with \\textit{Herschel}. % Our results provide important empirical support for a proposed prescription for the wavelength dependence of the cirrus noise, which incorporates a factor which varies directly as the SED.  In practice the noise will be best evaluated, most free of assumptions, simply by measuring $P_0^{1/2}$ at each wavelength and using equation~(\\ref{sigcir})."
    },
    "0910/0910.0821_arXiv.txt": {
        "abstract": "As a step toward a complete theoretical integration of 3D compressible hydrodynamic simulations into stellar evolution, convection at the surface and sub-surface layers of the Sun is re-examined, from a restricted point of view, in the language of mixing-length theory (MLT) .  Requiring that MLT use a hydrodynamically realistic dissipation length gives a new constraint on solar models. While the stellar structure which results is similar to that obtained by YREC \\citep{gdkp92,bp04} and Garching models \\citep{swl97}, the theoretical picture differs. A new quantitative connection is made between macro-turbulence, micro-turbulence, and the convective velocity scale at the photosphere, which has finite values.  The ``geometric parameter'' in MLT is found to correspond more reasonably with the size of the strong downward plumes which drive convection \\citep{sn98}, and thus has a physical interpretation even in MLT. Use of 3D simulations of both adiabatic convection and stellar atmospheres will allow the determination of the dissipation length and the geometric parameter (i.e., the entropy jump), {\\em with no astronomical calibration.}  A physically realistic treatment of convection in stellar evolution will require additional modifications beyond MLT, including effects of kinetic energy flux, entrainment (the most dramatic difference from MLT found by \\cite{ma07b} ), rotation,  and magnetic fields \\citep{balhawl,balbus}. ",
        "introduction": "Recent simulations of three-dimensional compressible convection and their theoretical analysis \\citep{ma07b,amy09} have shown that the interpretation of mixing-length theory (MLT), as currently used in stellar evolution \\citep{bv58,cox68,clay83,hansen,kippen} is flawed. This mixing length $\\ell$ is parameterized as $ \\aml = \\ell /H_P$, where $H_P$ is the local pressure scale height, and $\\aml$ is adjusted to reproduce the radius of the present-day Sun. However, instead of being an adjustable parameter, the mixing length is found to correspond to the dissipation length $\\ell_d$ of the turbulence \\citep{kolmg41,kolmg},  and determined by the size of the largest eddies \\citep{ma07b,amy09}. From our own simulations \\citep{ma09} and those of others we find a robust tendency for the dissipation length to be \\begin{equation} \\ell_d \\approx \\min ( \\ell_{CZ}  , 4 H_P ), \\label{eq-alpha} \\end{equation} where $\\ell_{CZ}$ is the depth of the convection zone. For shallow convection zones, the dissipation length is limited by the depth of the convective region, and  seems to approach a limiting value of $\\ell_d \\approx 4 H_P$ for deep convection zones.  If $\\aml$ is fixed, other parameters in MLT, which are generally left fixed by historical convention, may be adjusted to compensate (e.g., \\citep{tfw90,sc08}). The most significant of these parameters is the geometric factor\\footnote{This is essentially the $c$ factor of \\cite{tfw90}.}, which adjusts the rate at which radiation limits the degree of entropy excess in the super-adiabatic region (SAR). For simplicity we use $g_{ML}$ to denote the geometric factor in units of the value used in conventional MLT (see Appendix for details). If we identify the geometric parameter as a measure of the size of the SAR, we remove the last free parameter in MLT. Although MLT is an incomplete theory, it does serve as a useful \"language\" to explain some of the changes implied by  3D simulations.   The mixing length theory itself  \\citep{vitense53,bv58}, if used consistently, does capture many (but not all) aspects of turbulent convection. However, a real replacement for MLT will provide a global solution and relax the local connection between the superadiabatic gradient and the enthalpy flux, so that regions of the convective zone can be subadiabatic, as observed in simulations.  In order to establish a ``baseline'' from which to compare new effects demanded by numerical simulations and by laboratory experiment (e.g., fluctuations, non-locality, and entrainment), the framework of the standard solar model \\citep{bsp04} is examined with respect to modification of some aspects of convection. In this paper we show that the addition of dynamically realistic values of mixing length and geometric factor give some interesting insights into the nature of the average stratification of the Sun just below the photosphere. In Section~2 we construct a series of solar models with $\\aml$---$\\gml$ pairs, to delineate their properties. The notion of a ``standard solar model'' derives from the work of John Bahcall and collaborators, and is summarized in \\cite{bahcall}. It represents  what is probably the most carefully tested aspect of the theory of stellar evolution. In Section~3 we compare our models to standard results using the Yale Rotational Evolution Code \\citep{gdkp92,bp04}, and the Garching code (see \\cite{schlattl,swl97}). In Section~4 we compare the outer layers of our models to the 3D atmospheres of Nordlund and Stein \\citep{aspat}, semi-empirical models of the solar atmosphere \\citep{fontenla}, and re-examine the question of convective velocities at the photosphere. In Section~5 we summarize the implications of this work. ",
        "conclusions": "Insights from 3D compressible convection simulations and theory \\citep{ma07b,amy09} have been applied to sub-photospheric regions of solar models. Even within MLT, a dynamically consistent velocity field (i.e., a consistent choice of $\\aml$ and an adjusted $\\gml$), gives a better agreement with \\begin{enumerate} \\item empirical $T(\\tau)$ relations, and \\item 3D hydrodynamic models of stellar atmospheres. \\end{enumerate}  Using the correct condition for mixing (the bulk Richardson number) implies that the 1D atmospheres should exhibit hydrodynamic flow. Further, simple hydrodynamic considerations {\\citep{press81,pr81} suggest g-mode waves will be generated  and penetrate to the photosphere (these are generated by turbulent forcing from convection). We show that there is a connection between the predicted turbulent velocity scale and the observed (macro and micro)-turbulent velocities, which removes the embarrassment of non-convective surface regions predicted by 1D stellar atmosphere theory. As a bonus, we find that the observed macro- and micro-turbulence for the Sun can be used to fix the choice of $g_{ML}$ (model~D).  We may also have a resolution of an apparent contradiction.  Atmospheric models of white dwarfs \\citep{winget,bergeron}, which have shallow convection zones, use MLT parameters (ML2: $\\alpha = 0.6$, considerably smaller than used for the Sun), indicating less vigorous convection. Low mass eclipsing binaries \\citep{stassun,morales} are generally fit with $\\alpha \\sim 1$ (again low convective efficiency), these models do {\\em not} have shallow convection zones. In MLT there is no rationale for these differences. Use of Eq.~\\ref{eq-alpha} will give models having thin convection zones which agree with MLT models using small $\\alpha_{ML}$, so we expect to reproduce the white dwarf results. For low mass stars, the surface temperatures will be lower than the solar value, so that the SAR should comprise more mass, i.e., we expect larger $g_{ML}$ to be physically correct. Table~\\ref{tablep} indicates that there is a trade off between $\\alpha_{ML}$ and $g_{ML}$: to compensate for lower $g_{ML}$, $\\alpha$ must be lower, for the same radius. For a deep convection zone, $\\alpha_{ML}$ is fixed; then a stronger SAR (larger $g_{ML}$, and more inefficient convection) will give a larger radius. This is the sense of the discrepancy of the computed radii for low mass eclipsing binaries \\citep{stassun,morales}, and we suggest that part of the discrepancy may be due to the convection algorithm used. Unfortunately, direct calculation of low mass dwarfs ($M \\approx 0.2 \\rm M_\\odot$) with $\\alpha_{ML} \\approx 4$ exposes limitations in MLT: the SAR is forced upward into the photosphere, making 3D atmospheres a necessity for gaining insight into a plausible treatment in stellar models.    We have, in fact, sketched a way to eliminate astronomical calibration from stellar convection theory: \\begin{enumerate} \\item  Adjust $\\aml$ from convection simulations. The mixing length is $\\ell_m = \\aml\\   H_P$ (where $\\ H_P$ is the pressure scale height), and equal to the depth of the convection up to $4\\   H_P$, and $\\aml \\approx 4$ for deeper convection zones. \\item Adjust $\\gml$ from 3D hydrodynamic atmosphere simulations, fitting the curve of super-adiabatic excess in the superadiabatic region. This is more accurate, but equivalent to adjusting $g_{ML}$ to reproduce a self-consistent SAR. \\end{enumerate} Notice that a fit to the present day solar radius is not logically necessary.  By seriously considering MLT, we have determined that no significant free parameters are left to adjust within the framework of the theory. We find that the choice of two characteristic lengths, {\\em which are determined by the flow}, close the system: the turbulent dissipation length and the size of the super-adiabatic region (SAR). Alternatively, the constraint that the observed micro- and macro-turbulent velocities agree with those predicted using the bulk Richardson criterion for surface convective mixing can be used instead of the SAR size.  MLT is still an incomplete theory, but it is suggestive that even modest changes toward a better physical interpretation, based upon 3D simulations and on a more complete turbulence theory, do give improvements in the models. MLT, as used here, may be derived from a more general turbulent kinetic energy equation by ignoring certain terms \\citep{amy09}. Some of the ignored terms are important, emphasizing that MLT is incomplete. However, the approach sketched above may be generalized, and inclusion of missing terms gives a convection theory that is nonlocal, time dependent, provides robust velocity estimates, and is based on simulations and terrestrial experiment, with no astronomical calibration. This more difficult theory will be presented in detail in future publications."
    },
    "0910/0910.0010_arXiv.txt": {
        "abstract": "Transiting planet discoveries have yielded a plethora of information regarding the internal structure and atmospheres of extra-solar planets. These discoveries have been restricted to the low-periastron distance regime due to the bias inherent in the geometric transit probability. Monitoring known radial velocity planets at predicted transit times is a proven method of detecting transits, and presents an avenue through which to explore the mass-radius relationship of exoplanets in new regions of period/periastron space. Here we describe transit window calculations for known radial velocity planets, techniques for refining their transit ephemerides, target selection criteria, and observational methods for obtaining maximum coverage of transit windows. These methods are currently being implemented by the Transit Ephemeris Refinement and Monitoring Survey (TERMS). ",
        "introduction": "\\label{introduction}  Planet formation theories thus far extract much of their information from the known transiting exoplanets which are largely in the short-period regime. This is because transit surveys which have provided the bulk of the transiting planet discoveries, such as SuperWASP \\citep{pol06} and HATNet \\citep{bak02}, are biased towards this region. The two planets which contribute to the sample of intermediate to long-period transiting exoplanets are: HD~17156b \\citep{bar07a} and HD~80606b \\citep{lau09,mou09}, the latter of which exhibits both a primary transit and secondary eclipse. In both of these cases, the detection was largely due to the inflated transit/eclipse probability caused by their extreme orbital eccentricities. Both of these were observed to transit through predictions based upon their radial velocity data.  Planetary orbits may be considered in three basic categories: short-period ($< 10$ days) planets, intermediate to long-period high-eccentricity ($> 0.1$) planets, and intermediate to long-period low-eccentricity planets. The first type of planet are the focus of current studies, and we are beginning to gain insight into the second type (HD 80606b for example). Exploration into the structure of those planets occupying the second and third orbit types will require probing parameter-space beyond that currently encompassed by the known transiting exoplanets. The science objectives of such an exploration include understanding how planetary properties, such as average planet density, vary with periastron distance as well as providing the first observational data for exoplanet models with low incident stellar radiation. Recent observations of HD~80606b by \\citet{gil09} and \\citet{pon09} suggest a spin-orbit misalignment caused by a Kozai mechanism; a suggestion which could be investigated in terms of period and eccentricity dependencies if more long-period transiting planet were known.  Figure \\ref{peri} shows the distribution of periastron distances for the known exoplanets using data from Jean Schneider's Extra-solar Planets Encyclopaedia\\footnote{http://exoplanet.eu/}. Also shown as a shaded histogram is the same distribution for the known transiting exoplanets. Clearly the periastron distribution of the known transiting planets does not accurately represent the distribution of the entire sample of known exoplanets. Thus far, our picture of exoplanets is based on the planets with super-heated atmospheres through either short-period orbits or intermediate-period eccentric orbits. Figure 1 shows how our view of planetary properties is highly biased towards planets of short periastron distance, thus demonstrating the need for detecting transits of exoplanets at larger distances in order to study atmospheres with greatly reduced flux from the parent star.  \\begin{figure} \\includegraphics[width=8.2cm]{f1.eps} \\caption{Histogram of periastron distances of known exoplanets. The shaded region corresponds to known transiting planets, the population of which is comprised of planets in short-period orbits and intermediate-period eccentric orbits.} \\label{peri} \\end{figure}  Long-period transits will give us valuable insight into the structure of exoplanets that are more similar to those in our own solar system. Additionally, host stars of long-period planets discovered using the radial velocity technique tend to be very bright, and potentially one of these could be the brightest star with a transiting Jupiter-mass planet. The brightness of such host stars can faciliate atmospheric studies via transmission spectra \\citep{bur06,red08,sne08} and thermal phase variations \\citep{knu07}. The long period typically also translates to a longer transit duration, allowing significantly more signal to be collected when trying to probe the atmospheric composition of these objects with high-resolution spectrographs. Since most bright late-F,G,K type stars have already been searched for short-period Jupiter-mass planets, the long period hosts are the obvious sample to now search since their relative brightness and longer transit times are both very helpful in attempting to probe the exoplanetary atmosphere.  It is clear that the discovery of long-period transiting planets is very important. They are, however, difficult to detect in ground-based transit surveys due to the the presence of correlated (red) noise \\citep{pon06} and effects of the observational window function \\citep{von09}. Though combining data from various on-going transit surveys to search for transiting long-period planets will improve detectability \\citep{fle08}, it is easier to discover transits among the large number of long-period exoplanets already known from radial velocity surveys. Based upon Monte-Carlo simulations of the transit probabilities, it is expected that several of the known long-period planets should transit their parent stars due to the enhanced probabilities produced by eccentric orbits \\citep{bar07b,kan08a}. There have been suggestions regarding the strategy for photometric follow-up of these radial velocity planets at predicted times of primary transit \\citep{kan07a} and secondary eclipse \\citep{kan09a}, and the instruments which could be used for such surveys \\citep{lop06}. Here we describe a methodology through which to search for transits of known planets. In Section 2 we calculate transit windows for a large selection of the known exoplanets and show the impact of additional radial velocity measurements on refining the transit ephemerides. Section 3 describes the techniques through which optimal target selection and observations of the transit window can be achieved. The methods described here have been successfully tested and implemented by the Transit Ephemeris Refinement and Monitoring Survey (TERMS). ",
        "conclusions": "Many of the known radial velocity planets have yet to be surveyed for transit signatures. The detection of a transit for the intermediate to long-period planets would add enormously to our knowledge of planetary structure and, in particular, how the structure varies with semi-major axis and periastron distance. The advantages of targeting long-period radial velocity planets are the brightness of the host stars and the prior knowledge of the planetary orbital parameters. However, a major challenge of monitoring the host stars at predicted transit times is that many transit windows have deteriorated over time, such that the telescope time required renders attempts to do so impractical.  We have shown through calculations for 245 of the known exoplanets how the size of the transit window varies with period and geometric transit probability. The large uncertainties associated with the transit mid-point for the long-period planets is dominated by the uncertainties in the period and time of periastron passage estimated from the discovery data. We demonstrated, using the examples of HD~190228 and HD~231701, that a handful of carefully timed additional measurements can vastly improve the size of the transit window and thus bring the monitoring of the window into the reach of ground-based programs.  The difficulties involved in the observing schedule largely result from matching transit windows with the observability of the targets, particularly for long-period planets whose transit windows are widely spaced. We have described a planning strategy which will make optimal use of both pre-allocated and service telescope time, also noting the advantages of both longitude coverage and space-based observations. It is important to consider for the scheduling that the central part of the transit window can be significantly more valuable than the wings of the window, depending upon the nature of the orbital parameter uncertainties.  The described techniques and science goals are currently being undertaken and investigated by the Transit Ephemeris Refinement and Monitoring Survey (TERMS). Note that the observations from this survey will lead to improved exoplanet orbital parameters and ephemerides even without an eventual transit detection for a particular planet. The results from this survey will provide a complimentary dataset to the fainter magnitude range of the Kepler mission \\citep{bor09}, which is expected to discover many transiting planets including those of intermediate to long-period planets."
    },
    "0910/0910.1421_arXiv.txt": {
        "abstract": "{The Effelsberg-Bonn HI Survey (EBHIS) covers the whole sky north of Dec(2000)$\\, =\\, -5^{\\rm \\circ}$ on a fully sampled angular grid. Using state-of-the-art FPGA-spectrometers we perform a Milky Way and an extragalactic HI survey in parallel. Moreover, the high dynamic range and short dump time of the HI spectra allow to overcome the vast majority of all radio-frequency-interference (RFI) events. The Milky Way data will be corrected for the stray-radiation bias which warrants a main-beam efficiency of 99\\%. Towards the whole survey area we exceed the sensitivity limit of HIPASS, while towards the Sloan-Digital-Sky-Survey (SDSS) area EBHIS offers an order of magnitude higher mass sensitivity. The Milky Way data will be a cornerstone for multi-frequency astrophyics, while the extragalactic part will disclose detailed information on the structure formation of the local universe.}  \\FullConference{Panoramic Radio Astronomy: Wide-field 1-2 GHz research on galaxy evolution - PRA2009\\\\ June 02 - 05 2009\\\\ Groningen, the Netherlands}  \\begin{document} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.4803_arXiv.txt": {
        "abstract": "Within the interacting Fermi gas model for isospin asymmetric nuclear matter, effects of the in-medium three-body interaction and the two-body short-range tensor force due to the $\\rho$ meson exchange as well as the short-range nucleon correlation on the high-density behavior of the nuclear symmetry energy are demonstrated respectively in a transparent way. Possible physics origins of the extremely uncertain nuclear symmetry energy at supra-saturation densities are discussed. ",
        "introduction": "The density dependence of nuclear symmetry energy $E_{sym}(\\rho)$ is currently a key issue in both nuclear physics and astrophysics, see e.g., Refs.\\cite{li3,bro,li2,dan,bar,li1,Sum94,Bom01,Lat04,Ste05a,Lee09}. Despite much theoretical and experimental effort, our current knowledge about the $E_{sym}(\\rho)$ is still rather poor especially at supra-saturation densities. Experimentally, some constraints on the $E_{sym}(\\rho)$ at sub-saturation densities have been obtained recently from analyzing nuclear reaction data, see, e.g., Refs. \\cite{LWC05,tsa,Cen09}. At supra-saturation densities, however, the situation is much less clear because of the very limited data available and the few model analysis carried out so far although some indications of a super-soft $E_{sym}(\\rho)$ at high densities have been obtained from analyzing the $\\pi^+/\\pi^-$ ratio in relativistic heavy-ion collisions\\cite{xia}. Theoretically, essentially all available many-body theories using various interactions have been used in calculating the $E_{sym}(\\rho)$. Unfortunately, the predictions at supra-saturation densities are very diverse, for a recent review, see, e.g., Ref. \\cite{li1}. Assuming all models are equally physical and noting that there is no first principle guiding its high-density limit, it is fair to state that the $E_{sym}(\\rho)$ at supra-saturation densities is currently still completely undetermined. So, why is the $E_{sym}(\\rho)$ so uncertain at supra-saturation densities? This is obviously an important question that should be addressed timely especially since several dedicated experiments have now been planned to investigate the high density behavior of the $E_{sym}(\\rho)$ at CSR/China \\cite{CSR}, GSI/Germany \\cite{GSI}, MSU/USA \\cite{MSU} and RIKEN/Japan \\cite{RIKEN}. Identifying the causes for the uncertain high-density $E_{sym}(\\rho)$ may help experimentalists to decide what experiments to do and what observables to measure. While we can not fully answer this question, we identify several important factors and demonstrate their effects on the high-density behavior of the $E_{sym}(\\rho)$ using probably the simplest many-body theory available, namely the interacting Fermi gas model for isospin asymmetric nuclear matter, e.g., Ref. \\cite{pre}. There are many long-standing physical issues on how to treat quantum many-body problems at high densities, the various techniques used in different many-body theories may be among the possible origins of the very uncertain $E_{sym}(\\rho)$ at supra-saturation densities. Nevertheless, it is still very useful to examine effects of some common ingredients used in most many-body theories, such as the three-body and tensor forces, in the simplest model possible. While the interacting Fermi gas model can not be expected to describe all properties of infinite nuclear matter and finite nuclei as accurately as those more advanced microscopic many-body theories, it does give an analytical expression for the $E_{sym}(\\rho)$ in terms of the isospin-dependent strong nucleon-nucleon (NN) interaction in a physically very transparent way \\cite{xuli2}. Most importantly, the key underlying physics responsible for the uncertain $E_{sym}(\\rho)$ at supra-saturation densities can be clearly revealed. In particular, effects of the spin-isospin dependent effective three-body force, the density dependence of the in-medium short-range tensor forces and the short-range nucleon correlation can be demonstrated clearly. The results are expected to be useful for not only understanding predictions of the various many-body theories but also ultimately determining the $E_{sym}(\\rho)$ at supra-saturation densities. ",
        "conclusions": "In summary, the high density behavior of the symmetry energy $E_{sym}$ has long been regarded as the most uncertain property of dense neutron-rich nuclear matter because even its trend is still controversial. Within the interacting Fermi gas model, the $E_{sym}$ can be expressed in terms of the isospin-dependent NN strong interaction. The high-density behavior of the $E_{sym}$ is determined by the competition between the in-medium isosinglet and isotriplet nucleon-nucleon interactions. Respective effects of the in-medium three-body interaction and the short-range tensor force on the high-density behavior of the $E_{sym}$ are examined separately. It is found that the strength of the spin (isospin) dependence of the three-body force and the in-medium $\\rho$ meson mass in the short-range tensor force are the key parameters controlling the high-density behavior of the $E_{sym}$. These findings are useful for understanding why the nuclear symmetry energy is very uncertain at supra-saturation densities.  We would like to thank Lie-Wen Chen, Wei-Zhou Jiang, Che Ming Ko, Hyun Kyu Lee, Rupert Machleidt, Mannque Rho and Wei Zuo for helpful discussions. One of us (B. A. Li) would also like to thank the kind hospitality he received at the World Class University program at Hanyang University in Seoul, Korea where he benefited from the stimulating discussions on the main issues studied in this work. This work is supported by the US National Science Foundation Awards PHY-0652548 and PHY-0757839, the Research Corporation under Award No.7123 and the Texas Coordinating Board of Higher Education Award No.003565-0004-2007, the National Natural Science Foundation of China (Grants 10735010, 10775068, and 10805026) and by the Research Fund of Doctoral Point (RFDP), No. 20070284016."
    },
    "0910/0910.3048_arXiv.txt": {
        "abstract": "We investigate an asteroseismic model of non-rotating paramagnetic neutron star with core-crust stratification of interior pervaded by homogeneous internal and dipolar external magnetic field, presuming that neutron degenerate Fermi-matter of the star core is in the state of Pauli's paramagnetic permanent magnetization caused by polarization of spin magnetic moments of neutrons along the axis of magnetic field of collapsed massive progenitor. The magnetic cohesion between metal-like crust and permanent-magnet-like core is considered as playing a main part in the dynamics of starquake. This process is thought of as impulsive release of magnetic field stresses, brought about by disruption of magnetic field lines on the core-crust interface resulting in the crust fracture, which is followed by tectonic shear displacements of crust against massive core. Focus is laid on the post-quake relaxation of the star by node-free torsional vibrations of highly conducting crustal solid-state plasma, composed of nuclei embedded in the degenerate Fermi-gas of relativistic electrons, about axis of magnetic field frozen in the immobile paramagnetic core. Two scenarios of these axisymmetric seismic vibrations are examined, in first of which these are considered as maintained by combined action of Lorentz magnetic and Hooke's elastic forces and in second one by solely Lorentz force. The damping is attributed to shear viscosity of crustal material. Based on the energy variational method of magneto-solid-mechanical theory of elastic continuous medium, the spectral formulae for the frequency and lifetime of this toroidal mode are obtained and discussed in the context of theoretical treatment of recently discovered quasi-periodic oscillations of the X-ray outburst flux from SGR 1806-20 and SGR 1900+14 as being produced by above seismic vibrations. ",
        "introduction": "Studying seismic vibrations of neutron stars is among the most promising tools in search for dynamical laws governing continuum mechanics and macroscopic electrodynamics of extremely dense matter compressed by self-gravity to the nearly normal nuclear density at which microscopic composition of stellar material is dominated by neutron component. One of the most  conspicuous features distinguishing this subclass of final stage (FS) stars of evolutionary track from  their massive main-sequence (MS) progenitors is that they come into existence as sources of ultra powerful magnetic fields (Chanmugam 1992). It is hoped, therefore, that exploring vibrational behavior of spherical material masses capable of accommodating super strong magnetic fields we can gain some insight into the nature of the nuclear matter magnetism.  In this connection it is appropriate to emphasize that contrary to liquid MS stars whose magnetic fields are generated in dynamo processes by flows of perfectly conducting flowing matter (energy supply of persistent convection in which comes from cental thermonuclear reactive zones), in the FS stars, white dwarfs and pulsars, there are no nuclear energy sources to  support persistent convection (e.g. Flowers \\& Ruderman 1977, Spruit 2008). Such point of view invalidates applicability to FS stars of magneto-fluid-mechanical arguments and methods in the analysis of magnetic field effects such, for instance, as anti-dynamo theorems. Unlike the dynamo-generated fairly weak fields of liquid MS stars, the super strong and highly stable to spontaneous decay magnetic fields of isolated neutron stars (Bhattacharya \\& van den Heuvel 1991, van den Heuvel 1994, Bhattacharya 2002) are regarded as fossil (e.g. Ferrario and Wickramasinghe 2004). These are thought of as frozen in the immobile, non-convective, degenerate stellar matter possessing properties of elastically deformable solid. It is believed, therefore, that basic features of asteroseimology of strongly magnetized end-products of stellar evolution (white dwarf stars, neutron stars and still hypothetical quark stars) can properly be understood working from equations of solid-mechanical theory of continuous media\\footnote{It may be appropriate to mention here investigations in nuclear physics on theoretical treatment of giant resonances in terms of vibrations of ultra fine piece of continuous nuclear matter in which it has been found that behavior of atomic nucleus in these fundamental modes of nuclear response bears strong resembles to spheroidal and torsional vibrations of elastically deformable solid sphere rather than drop of flowing liquid. (e.g. Bastrukov et al 2008b and references therein). One of the remarkable analytic inferences of these investigations is that the shear modulus of nuclear material (which is treated, therefore, as an elastic Fermi-solid) entering equations of nuclear solid-globe model, is found to be precisely coinciding with the pressure of degenerate Fermi-system of nucleons. This finding of nuclear physics suggests that neutron star can be thought of as a spherical mass of an elastic Fermi-solid of the normal nuclear density in which the energy of gravitational pull is counterbalanced by elastic energy stored in the incessant Fermi-motion of neutron quasi-particles in the potential of self-gravity (e.g. Bastrukov et al 2007a). Bearing this in mind and widely spread opinion that a neutron star can be regarded as an astrophysical counterpart of an atomic nucleus (e.g. Lipunov 1992, Bisnovatyi-Kogan 2002), it has been shown in (Bastrukov et al 1999a, 1999b) that equations of nuclear elastodynamics, originally introduced in the nuclear physics of giant resonances as fundamental equations of continuum mechanics of nuclear matter, can be successfully utilized for the study of neutron star pulsations. And in recent work (Bastrukov et al 2007a), it was shown that linearized equations of nuclear elastodynamics and canonical equations of linear theory of elasticity lead to identical results regarding frequencies of fundamental modes of node-free spheroidal and torsional vibrations of spherical mass of viscoelastic solid which are of particular interest for our present discussion.} (McDermott et al 1988, Blaes et al 1989, Strohmayer 1991, Bastrukov 1996, Bastrukov et al 1999a, 1999b, Franco et al 2000, Bastrukov et al 2007 see, also, references therein).  Many agree today that recent observations of quasi-periodic oscillations (QPOs) in the outburst X-ray flux of two soft gamma-ray repeaters, SGR 1806-20 and SGR 1900+14 (Israel et al 2005, Watts \\& Strohmayer 2006) have offered a real opportunity of using the detected QPO frequencies for asteroseismological inferences about internal structure, mechanical and magnetic properties of material of these ultra magnetized neutron stars whose X-ray bursting activity exhibits many features in common with earth quakes (Chen et al 1996, Woods \\& Thompson 2006, Mereghetti 2008). The detected QPOs frequencies (measured in Hz), which are of particular interest for the subject of present work, are \\begin{eqnarray} \\label{e1.1} && \\mbox{\\rm SGR}\\, 1806-20: 18, 26, 29, 92, 150, 625, 1840;\\\\ \\label{e1.2} && \\mbox{\\rm SGR}\\, 1900+14:  28, 54, 84, 155. \\end{eqnarray} With all above in mind, in (Bastrukov et al 2007-2009) several models of post-quake vibrational relaxation of a solid star have been investigated. Particular attention has been given to the regime of node-free or nodeless shear axisymmetric vibrations which remain less investigated in  theoretical seismology as compared to the standing-wave regime of vibrations of both solid FS stars and such solid celestial objects as Earth-like planets (e.g. Lapwood \\& Usami 1981, Lay \\& Wallace 1995, Aki \\& Richards 2003). In particular, in (Bastrukov et al 2007b, 2008a), a case of the elastic-force-driven nodeless shear oscillations, both torsional --  $_0t_\\ell$ and spheroidal -- $_0s_\\ell$, entrapped in the finite-depth crust $\\Delta R$ of core-crust model of a solid star has been studied in some details with remarkable inference that dipole overtones of spheroidal and torsion nodeless elastic shear vibrations of crust against immobile core exhibit features generic to Goldstone soft modes. On the other hand one can casts doubt on arguments of model presuming that detected QPOs have been brought about by vibrations restored by solely solid-mechanical force of Hooke's elastic stresses because such interpretation rests on the poorly justifiable assumption of the dynamically passive role of ultra strong magnetic field, as defining only direction of the axis about which the stellar matter undergo differentially rotational fluctuations, i.e., that magnetic field frozen in the star remains unaltered in this process. Bearing this in mind and assuming that the admixture of charged particles (basically electrons and protons) imparts to neutron-dominated Fermi-matter of these star mechanical properties of perfectly conducting material, it seems more plausible to expect that post-quake relaxation should be strongly affected by Lorentz force of magnetic field stresses setting highly conducting material in rotational motion about axis of the dipole magnetic moment of the star.  With this in mind, in (Bastrukov et al 2009a, 2009b), the above data have been analyzed in the model of a solid star undergoing global torsional nodeless vibrations about magnetic axis under the joint action of Lorentz force of magentic field stresses and  Hooke's force of elastic shear stresses. And it was found that such a model provides fairly reasonable account of general trends in data on QPO frequencies for SGR 1900+14 and data for SGR 1900+14 for frequencies from the range $30\\leq \\nu \\leq 200$ Hz, but faces serious difficulties in interpreting the low-frequency vibrations with $\\nu=18$ and $\\nu=26$ Hz in the SGR 1806-20 data. Also, the model of global torsional vibrations of solid star with froze-in homogeneous magnetic field leaves some uncertainties regarding the nature of vibrations with $\\nu=625$ and  $\\nu=1840$ Hz. This last issue has been scrutinized in recent work (Bastrukov et al 2009b) from the standpoint of a solid star model with non-homogeneous poloidal magnetic field of well-known Ferraro's form and was found that these high-frequency QPOs can be properly interpreted as being produced by global Alfv\\'en node-free torsional vibrations of very high overtones.  In the meantime, it has been argued some time ago in the studies of pulsar glitches as being produced by star quakes (Bastrukov et al 1996, 1997, 1999b) that post-glitch relaxation should most likely be dominated by Alfv\\'en oscillations of crustal electron-nuclear solid-state plasma (composed of immobile nuclei dispersed in degenerate Fermi-gas of relativistic electrons) about axis of magnetic field frozen in the immobile core (and, also, by less dense gaseous plasma expelled from the surface by starquake). From the view point of core-crust model of quaking neutron star (Franco et al 2000),  there are two way of thinking why only peripheral region of the star can be set in Lorentz-force-driven Alfv\\'en vibrations.  First belong to a case when strength of seismic tremor is fairly week and perturbation sets in vibrations only finite-depth crustal region of the star, leaving massive core unaltered. The second way, when core material just incapable of sustaining Alfv\\'en vibrations. One of the main conditions for existing of such vibrations is the extremely large (effectively infinite) electrical conductivity of stellar substance (e.g., Chandrasekhar 1961, Mestel 1999), but this may not be the case for core material whose microscopic composition is dominated, according to canonical understanding of neutron star, by matter possessing properties of non-conducting degenerate Fermi-gas of non-relativistic neutrons.   \\begin{figure} \\centering{\\includegraphics[width=8.0cm]{f1.eps}} \\caption{ The internal constitution of two-component, core-crust, model of paramagnetic neutron star. The massive core is considered as a poorly conducting permanent magnet composed of degenerate Fermi-gas of non-relativistic neutrons in the state of Pauli's paramagnetic saturation caused by field-induced alignment of spin magnetic moment of neutrons along the axis of uniform internal and dipolar external magnetic field frozen in the star on the stage of gravitational collapse of its MS progenitor. A highly conducting metal-like material of the neutron star crust, composed of nuclei embedded in the super dense degenerate Fermi-gas  of relativistic electrons, is regarded as  electron-nuclear solid-state magneto-active plasma capable of sustaining Alfv\\'en oscillations.} \\end{figure}    The physically meaningful framework for discussion of this second case provides the two-component, core-crust, model of quaking paramagnetic neutron star, pictured in Fig.1, that has been extensively discussed some time ago in works \\footnote{The main purpose of these works was to explore potential capabilities of equations of magneto-elastic dynamics (MED) suggested as fundamental equations of continuum mechanics and macroscopic electrodynamics of non-conducting but highly magnetically polarized stellar matter that have been formulated in a manner of equations of magnetohydrodynamics underlying continuum-mechanical treatment of motions of perfectly conducting liquid-state and solid-state plasmas pervaded by magnetic fields. It was found on the basis of MED equations that magnetically polarized paramagnetic neutron matter, with spin-magnetic moments of neutrons aligned along frozen-in magnetic field, can transmit perturbation by torsional magneto-elastic wave having for such a matter one and the same physical meaning as Alfv\\'en wave does for magneto-active liquid-state or solid-state plasmas. Also, it was shown that a neutron star whose material is in the state of paramagnetic permanent magnetization can sustain solely torsional oscillations driven by force of antisymmetric magneto-polarizational stresses.} (Bastrukov et al 2001, 2002a, 2002b, 2003). In this model, the role of major source of super strong magnetic field of the star is attributed to permanently magnetized massive core composed of degenerate neutron matter in the the state of Pauli's paramagnetic saturation in which magnetic moments of neutrons turn out aligned in the direction of frozen-in magnetic field that has been inherited by neutron star from massive progenitor and amplified in the implosive process of its gravitational collapse. Unlike the non-conducting core matter, the material of crust matter resembles metallic solid. This suggests that it is magnetic cohesion between metal-like crust and permanent-magnet-like core, mediated by lines of ultra strong magnetic field, should play main part in the the starquake dynamics. In this paper we adopt this latter point of view by considering starquake as impulsive release of magnetic field energy by means of disruption of magnetic field lines on the core-crust interface and fracturing the crust by relived magnetic stresses. In some sense the magnetic field lines can be thought of as a super-hard piles endowing core-crust composition of neutron star with supplementary (to gravity forces) stiffness which is primarily responsible for its seismic stability to quake-induced tectonic shear displacements of crust against core\\footnote{The internal constitution of paramagnetic neutron star model under consideration does not exclude the existence of a fairly thin inner crust with average density several times less than the normal nuclear density (neutron drip-line) at which micro-composition of matter is dominated, as we know from nuclear physics, by paired neutrons. It must be stated clear, however, that such quasi-boson matter composed of paired neutrons has nothing to do with fairly dilute and flowing quasi-boson matter like superfluid Fermi-liquid composed of $^3$He atoms. According to current theory, the emergence of superfluidity in flowing boson material substances belongs to phase transition of macroscopic fraction of bosons to the state of Bose-Einstein condensation (BEC). One of the most striking features of such phase transition is that at the temperature of condensation the internal pressure of BEC state vanishes. The fact that Bose-matter composed of paired neutrons in BEC state is unable to withstand the pressure of self-gravity (contrary to the degenerate neutron Fermi-matter whose zero-temperature pressure of degeneracy counterbalancing the self-gravity pressure is central to the very notion of neutron star) suggests that such a matter, if exists, can only insufficiently contribute to the total mass budget of neutron star, otherwise the very notion of this compact object is loosing its original meaning. In the starquake of paramagnetic neutron star model under consideration, this inner crust material is thought of as operating like a lubricant facilitating non-compressional differentially rotational shear displacements of crust relative to much denser matter of massive core.}. It is expected that quake-induced disruptions of magnetic field lines, along which the outgoing radiation is developed, should be observed as glitches in the quiescent high-energy emission. These glitches are of substantially magnetic nature in the sense that they owe its origin to well-known effect of sudden, jump-like, increase in the magnetic field strength triggered by thermonuclear explosion. With this in mind we focus on reaction of the star on above quake-induced perturbation by node-free differentially rotational, torsional, vibrations of crustal solid-state plasma about axis of magnetic field frozen in the immobile paramagnetic core.    In section 2, a brief outline is given of the equations of solid-magnetics appropriate for the perfectly conducting viscoelastic continuous medium pervaded by a magnetic field. In section 3, the spectral formulae for the frequency and lifetime of this toroidal mode of non-radial nodeless vibrations driven by combined magnetic Lorentz and elastic Hooke's forces are obtained. The relevance of considered asteroseismic model to the detected QPOs in the X-ray flux during the flare of SGR 1806-20 and SGR 1900+14 is assessed in section 4. The newly obtained results are briefly summarized in section 5. ",
        "conclusions": "The neutron stars, both pulsars and magnetars, are seismically active compact objects which are formed, as is widely believed, in the magnetic-flux-conserving supernova collapse of the cores of massive stars coming into existence as sources of super strong magnetic field with a fairly complex internal constitution. The internal micro-composition of these compact objects has been and still remains among the most active investigations on search for adequate equations of state of neutron star matter (e.g. Weber 1999, Lattimer \\& Prakash 2007, Sedrakian 2007, Haensel, Potekhin \\& Yakovlev 2007).  Together with this development, the past three decades have seen increasing recognition that important insight into interior of neutron stars and properties of super dense and ultra magnetized matter provides asteroseimology which seeks to explain fast variability of their luminosity as being produced by seismic vibrations. Adhering to this attitude, we have investigated seismic response of core-crust model of quaking paramagnetic neutron star, pictured in Fig.1, by node-free torsional crust-against-core vibrations. In this model the source of ultra powerful and highly stable to spontaneous decay magnetic field is associated with  massive core, which is thought of as an ultra powerful spherical magnet. The physical nature of seismic stability of core-crust coupling is attributed to magnetic cohesion between core and metallic like crust\\footnote{At this point it may be appropriate to note that considered model permits co-existence of two sources and two compound parts of neutron star magnetism (current-carrying flows in the crust and field-induced spin paramagnetism of core) and consistent with the idea about strengthening of magnetic field intensity in young neutron stars, like magnetars, due to thermoelectric effect in the crust (Urpin \\& Yakovlev 1980, Blandford et al 1983). And also with the idea that damping of crustal currents, thought of as second source of magnetic field of pulsars and magnetars, serves as one of the main factors responsible for the decay of magnetic field of neutron stars (e.g. Goldreich  \\& Reisenegger 1992, Pons \\& Geppert 2008). It is expected that magnetic field decay should manifest itself in the lengthening of periods of Alfv\\'en oscillations.}. The fact that impulsive release of magnetic energy in the starquake, associated with disruption of magnetic field lines on the core-crust interface followed by fracturing the crust by relived magnetic stresses, suggests that model under consideration is in line with current understanding of X-ray bursting activity of soft gamma-ray repeaters. With this in mind, the obtained spectral formulae have been applied to asteroseismic analysis of data on quasi-periodic oscillations of the outburst X-ray emission from SGR 1900+14 and SGR 1806-20. In so doing we have investigated two scenarios of post-quake relaxation of paramagnetic neutron star by torsional vibrations of crust against immobile massive core distinguishing restoring forces of vibrations. In first case, the analysis of data on QPOs frequencies has been based on assumption that detected QPOs owe their existence to torsional nodeless vibrations restored by joint action of Lorentz magnetic and Hooke's elastic forces and we have found that obtained three-parametric spectral formula provides much better fit of data than two-parametric frequency spectrum of the early considered model of global magneto-elastic vibrations (Bastrukov et al 2009a). In second case of vibrations restored by solely Lorentz force of magnetic field stresses we found that obtained two-parametric frequency spectrum can too be fairly reasonably reconciled with detected QPOs frequencies. With all above in mind we conclude that main part in quake-induced torsional vibrations of crustal solid-state plasma about axis of magnetic field frozen in the immobile core plays Lorentz restoring force of magnetic field stresses and that observed frequencies of QPOs can be interpreted as an indirect evidence for permanent magnetization of super dense matter in the deep interior of magnetars.    The authors are grateful to Dima Podgainy (JINR, Dubna) for helpful assistance, Kinwah Wu for discission and Ziri Younsi (Mullard Space Science Laboratory, University College London) for critical reading of the manuscript on the initial stage of this work. This work is a part of projects on investigation of variability of high-energy emission from compact sources supported by NSC of Taiwan, grant numbers  NSC-96-2628-M-007-012-MY3 and NSC-97-2811-M-007-003."
    },
    "0910/0910.0757_arXiv.txt": {
        "abstract": "In this series of papers, we propose a theory to explain the formation and properties of rings and spirals in barred galaxies. The building blocks of these structures are orbits guided by the manifolds emanating from the unstable Lagrangian points located near the ends of the bar. In this paper we focus on a comparison of the morphology of observed and of theoretical spirals and rings and we also give some predictions for further comparisons. Our theory can account for spirals as well as both inner and outer rings. The model outer rings have the observed $R_1$, $R_1'$, $R_2$, $R_2'$ and  $R_1R_2$ morphologies, including the dimples near the direction of the bar major axis. We explain why the vast majority of spirals in barred galaxies are two armed and trailing, and discuss what it would take for higher multiplicity arms to form. We show that the shapes of observed and theoretical spirals agree and we predict that stronger non-axisymmetric forcings at and somewhat beyond corotation will drive more open spirals. We compare the ratio of ring diameters in theory and in observations and predict that more elliptical rings will correspond to stronger forcings. We find that the model potential may influence strongly the numerical values of these ratios. ",
        "introduction": "\\label{sec:intro}  The discs of barred galaxies very often have spectacular sub-structures such as grand design spirals, or inner and outer rings. In fact, these sub-structures are observed much more often in barred than in non-barred disc galaxies. Their formation is intimately linked to the evolution of disc galaxies and also gives important information on the corresponding underlying potential and the bar pattern speed. Therefore many efforts, both theoretical and numerical, have been made to explain them (see e.g. \\citealt{Lin67}, \\citealt{Toomre77}, \\citealt{Toomre81}, \\citealt{Atha84}, \\citealt{Atha.Bosma.85} and \\citealt{Buta.Combes96} for reviews). None of these theories, however, could explain all observed properties and, particularly, none could explain both spirals and rings.  For this reason, we propose in this and in previous papers a new theory to explain all these features in a single theoretical framework. It relies on the dynamics of the unstable Lagrangian points located near the ends of the bar and of the associated manifolds. These guide the motion of the chaotic orbits that constitute the building blocks of the spirals and rings. Our theory and its applications are described in five papers. The first two (\\citealt{RomeroGMAG06}, hereafter Paper I; \\citealt{RomeroGAMG07}, hereafter Paper II) contain most of the theoretical groundwork. They describe in detail the dynamics of the region around the two unstable Lagrangian points and of the corresponding manifolds. A theory, however, can be dynamically correct but still irrelevant to a particular application. Thus, the aim of the remaining three papers is to check whether our theory is applicable to observed spirals and rings. In Paper III \\citep{AthaRGM09}, we considered the orbits guided by the manifolds as potential building blocks for spiral and rings and derived their relevant properties. In this paper, the fourth of this series, we compare the morphology of observed and of theoretical spirals and rings to test whether there is agreement. We also make a number of predictions based on our theory, which can be tested with future observations. In Paper V (Athanassoula et al 2009b, to be submitted), the last of this series, we will present further comparisons with observations. These last three papers can be considered as one entity and therefore have the same main title (``Rings and spirals in barred galaxies''), but we will always consider them as part of the whole series and refer to them as Papers III, IV and V, respectively.  The comparisons between our theory and the observations have, by necessity, to be qualitative. Indeed, in all our calculations we have used three specific families of analytic potentials. These can not cover all possible forms of observed barred galaxy potentials, nor are we sure that they share all their relevant properties. Furthermore, we do not take into account  self-gravity, i.e. we neglect the potential of the structure itself. If a given quantity depends considerably on the properties of the potential, we may not find good quantitative agreement between our theoretical results and observations, even though our theory is correct and applicable. For example, if the ring extent depends strongly on the potential used, then one can expect discrepancies between the observed and the theoretical extent values, if the adopted potential does not match observations in the relevant regions.  Section \\ref{sec:theory} contains a brief reminder of the main aspects of our theory. This is necessarily simplified as well as condensed, so we refer to an interested reader to Papers I, II and III and to \\citealt{RomeroGMGA09} for more information, a more thorough description of the theory and references to other related work. Nevertheless, the main theoretical concepts and definitions are given so as to make this paper as stand-alone as possible, while avoiding repetition. In particular we introduce here the Lagrangian points, the invariant manifolds and the properties of the associated orbits. Section~\\ref{sec:morpho} compares the morphological types of our theoretically predicted density enhancements with the observed ones. Section~\\ref{sec:spiralprop} focuses on the properties of the theoretical and the observed spirals, such as their sense of winding, their number of arms and their pitch angle. Section~\\ref{sec:ringprop} is devoted to ring properties, such as the frequency of the various types of rings, the axial ratios of inner and of outer rings and the ratio of outer to inner ring major axes. Further comparisons with observations, including pattern speed prediction, the bar shape, kinematics, radial drift and abundance gradients will be discussed in Paper V, the last of this series. There we also give a global discussion of all the results in this series. For this reason we summarise only briefly the results of this paper in Sect.~\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  In previous papers (Papers I, II and III) we proposed a theory to explain the formation and properties of rings and spirals in barred galaxies. In this paper we compared their morphological properties to that of observed ones. Our theory can produce both spirals and inner and outer rings. Furthermore we get all three observed types of outer rings , i.e. $R_1$, $R_2$ and $R_1R_2$ rings and pseudo-rings, and no other morphologies. This means that all outer rings produced  by our theory have a qualitatively realistic morphology and that all observed morphologies are qualitatively reproduced. We also reproduce the dimples often seen clearly in $R_1$ rings. Furthermore, we predict that if the non-axisymmetric forcing is relatively low the resulting morphology will be an $R_1$ ring, while if it is stronger, it will be a spiral.  Spirals due to this theory start off from a location not far from the end of the bar and are trailing, in good agreement with observations. They are in the vast majority of cases two armed, although it is possible, if specific potentials are used, to have arms of higher multiplicities. Their shape is also realistic and we made a successful quantitative comparison with the arms of NGC 1365. They are able to reproduce the ``fall back'' of the arm towards the more inner region, seen in NGC 1365 and in many other barred spirals, something which density wave theory could not achieve. We also predict that stronger non-axisymmetric forces in the region around and somewhat beyond corotation will result in more open spirals.  Quantitative comparisons are less straightforward to make, not because they imply tighter constraints, which they do not, but because they involve not only the theory, but also the potential to which it is applied. We are confident that our theory is dynamically correct, and the comparisons with observations made here and in Paper V make us optimistic about its applicability to galaxies with spirals or rings. On the other hand, we are aware of the shortcomings of our potentials. First and foremost, our analysis is not self-consistent, so that we neglect the self-potential of any structure they create. Furthermore, the Ferrers' potentials, although the best of all the available analytical bar potentials, have a number of shortcomings, the most important one for our particular application being that the force they create diminishes too rapidly at large radii. The other potentials we used, although they lack this particular drawback, are ad hoc, and therefore not necessarily realistic. The above remarks should be kept in mind when assessing our quantitative comparisons. Our aim in such cases was to find whether the potentials we used give results in agreement with observations, or, in case of disagreement, whether any obvious and straightforward changes of the potential can bring agreement.  For the outer ring our analysis concentrates on $R_1$ cases, for which we have sufficient models to reach safe conclusions. With all the potentials used so far, the outer ring shape is in good agreement with observations. Most models are within one $\\sigma$ from the observed mean and all are within two $\\sigma$. We tried different changes of the potential -- such as different rotation curves, adding a lens, or stabilising the $L_1$ and $L_2$ Lagrangian points (see Paper III) -- and this result remained unchanged. There is thus excellent agreement between model and observed outer ring shapes.  For the inner ring, we found several models with inner ring shapes more elongated than observed. Examining by eye these cases, we found that for most, if not all, the inner manifolds correspond more to the outer regions of the bar and not to the inner ring. We can thus conclude that there is a fair agreement between theory and observations regarding the inner ring axial ratio.  Finally, the ratio of outer to inner ring major axes gives good agreement for some types of models, but not for all. Thus all pseudo-ring cases ($rR_1'$), independent of the model potential, give good agreement. This is the case also for both $rR_1$ and $rR_1'$ morphologies in models with a circular velocity curve showing a slight decrease at or somewhat beyond corotation, for models with stable $L_1$ and $L_2$ Lagrangian points and for a number of other models. On the other hand, agreement for cases with an $rR_1$ morphology and A or AE rotation curves (with or without a lens) often have too small a value of this ratio. It is thus clear that the underlying potential plays an important role. Our calculations suggest that models such that $Q_{t,L_1}$ is somewhat decreased within corotation and increased outside it would give good agreement for all ring size measurements, namely the inner and outer ring shapes (ratio of minor to major axes) and the ratio of the outer to inner ring major axes. It is, however, beyond the scope of this paper to find a model achieving this at all relevant radii, thus includes the best parts of all the models considered, particularly since this exercise may prove futile if self-consistency is taken into account.  We find a trend between the ring axial ratio and the relative amplitude of the non-axisymmetric forcing in the region around, or immediately beyond corotation, in the sense that stronger bars are linked with more elongated rings. This is true for both inner and outer rings. There is also a trend between the ratio of outer to inner ring major axes and this measure of non-axisymmetric forcing. If we focus on only one model family, or a couple of them, these trends can become clear correlations, but including all models broadens them, sometimes substantially.  In Paper V, the next and last paper of this series, we will present a number of other comparisons with observations and predictions, including kinematics, bar shape, radial drift and abundance gradients. We will discuss pattern speed prediction and give a global appreciation of the application of our theory to observed spirals and rings. We will also consider the effects of time evolution and of self-consistency. Finally we will discuss other theories and simulations and how they link to this work."
    },
    "0910/0910.2614_arXiv.txt": {
        "abstract": "{% Most solar and stellar dynamo models use the $\\alpha\\Omega$ scenario where the magnetic field is generated by the interplay between differential rotation (the $\\Omega$ effect) and a mean electromotive force due to helical turbulent convection flows (the $\\alpha$ effect). There are, however, turbulent dynamo mechnisms that may complement the $\\alpha$ effect or may be an alternative to it. } {% We investigate models of solar-type dynamos where the $\\alpha$ effect is completely replaced by two other turbulent dynamo mechanisms, namely the $\\vec{\\Omega}\\times\\vec{J}$ effect and the shear-current effect, which both result from an inhomogeneity of the mean magnetic field.} {% We studied axisymmetric mean-field dynamo models containing differential rotation, the $\\vec{\\Omega}\\times\\vec{J}$ and shear-current effects, and a meridional circulation. The model calculations were carried out using the rotation profile of the Sun as obtained from helioseismic measurements and radial profiles of other quantities according to a standard model of the solar interior. } {% Without meridional flow, no satisfactory agreement of the models with the solar observations can be obtained. With a sufficiently strong meridional circulation included, however, the main properties of the large-scale solar magnetic field, namely, its oscillatory behavior, its latitudinal drift towards the equator within each half cycle,  and its dipolar parity with respect to the equatorial plane, are correctly reproduced. } {% We have thereby constructed the first mean-field models of solar-type dynamos that do not use the $\\alpha$ effect. } ",
        "introduction": "The standard dynamo model for the Sun and stars is the $\\alpha\\Omega$ model where, within the framework of mean-field magnetohydrodynamics, the magnetic field is produced by an interplay between differential rotation (the $\\Omega$ effect) and the collective action of turbulent cyclonic convection flows, known as the $\\alpha$ effect \\citep{par55,par79,stk66,krarad80}. The $\\alpha$ effect is here responsible for generating the poloidal component of the large-scale magnetic field (LSMF), whose toroidal component is mainly generated by the the $\\Omega$ effect. The model is often supplemented with meridional flows, leading to so-called flux-transport dynamos \\citep[e.g.,][]{choschusdik95,kukrudschul01,rem06,dikgil07}. The meridional flows may transport toroidal magnetic flux toward the equator and their speed may determine the cycle period, thus allowing us to bypass a number of problems connected with the $\\alpha$ effect and $\\alpha\\Omega$ dynamos, as, for instance, that, in the case of the Sun, the obtained cycle periods are generally too short and  the magnetic activity is not sufficiently concentrated at low latitudes \\citep[see, e.g.,][]{oss03,rudhol04,brasub05}.  In mean-field magnetohydrodynamics, the influence of the turbulence on the LSMF is expressed by the mean turbulent electromotive force (MEMF), $\\vec{\\mathcal{E}}=\\left\\langle \\vec{u}\\times\\vec{b}\\right\\rangle $, where $\\vec{u}$ and $\\vec{b}$ are the fluctuating parts of the velocity and magnetic field and angular brackets denote averages. The by far best known contribution to $\\vec{\\mathcal{E}}$ is provided by the $\\alpha$ effect, namely, a turbulent electromotive force $\\alpha\\langle\\vec{B}\\rangle$, with $\\alpha$ denoting a {(symmetric)} tensorial factor of proportionality and  $\\langle\\vec{B}\\rangle$ the LSMF. However, there are other turbulent dynamo mechanisms besides the $\\alpha$ effect. Two of them are the $\\vec{\\Omega}\\times\\vec{J}$ effect \\citep{rad69,sti76} and the shear-current or $\\vec{W}\\times\\vec{J}$ effect \\citep{rogkle03,rogkle04}; $\\vec{\\Omega}$ is here the angular velocity of the stellar rotation, $\\vec{J}=\\nabla\\times\\langle\\vec{B}\\rangle/\\mu_0$ the large-scale electric-current density, and $\\vec{W}=\\nabla\\times\\vec{V}$ the large-scale vorticity, $\\vec{V}$ denoting the large-scale velocity. Both these effects result from an inhomogeneity of the LSMF, in contrast to the $\\alpha$ effect, which also works with a homogeneous $\\langle\\vec{B}\\rangle$ (that is to say, for calculating the $\\alpha$ effect, $\\langle\\vec{B}\\rangle$ may be considered as homogeneous on the scale of the fluctuations).  In the commonly used representation of the MEMF on the basis of symmetry considerations \\citep[see][]{rad80,krarad80,rad00,radklerog03}, the $\\vec{\\Omega}\\times\\vec{J}$ and shear-current effects represent contributions to the $\\vec{\\delta}$ term, {a term of the form $\\vec{\\mathcal{E}}_\\delta=\\vec{\\delta}\\times(\\nabla\\times\\langle\\vec{B}\\rangle)$, where $\\vec{\\delta}$ is a vector. Since $\\vec{\\mathcal{E}}_\\delta\\cdot\\vec{J}=0$, the effects described by this term cannot bring energy into the mean magnetic field and, thus, cannot lead to working dynamos when acting alone. } These effects have been investigated little in the context of solar and stellar dynamos. For a recent study of the possible role of the $\\vec{\\Omega}\\times\\vec{J}$ effect when acting together with the $\\alpha$ effect and differential rotation in a  spherical shell, {or when acting together with another part of the MEMF, not included in dynamo studies before,} in a rigidly rotation full sphere, see \\citet{pipsee09}, where an illustration of the physical mechanism behind the $\\vec{\\Omega}\\times\\vec{J}$ effect also may be found; the mechanism of the shear-current effect is very similar to that of the $\\vec{\\Omega}\\times\\vec{J}$ effect.  In this paper, we consider mean-field dynamo models in the geometry of a spherical shell, as appropriate for solar-type stars, where the $\\alpha$ effect is completely omitted. Instead, the $\\vec{\\Omega}\\times\\vec{J}$ and shear-current effects serve as turbulent dynamo mechanisms. In nearly all mean-field dynamo studies, the effective strengths of the different physical ingredients are controlled by freely varied dimensionless parameters; in the case of dynamo effects, e.g., the $\\alpha$ effect, these are usually referred to as dynamo numbers. This reflects our present knowledge of the physical processes in the convection zones of the Sun and stars. Realistic self-consistent numerical models of these processes and their interactions will remain out of reach for the foreseeable future. Given this situation, we deem it advisable to explore the potentials of turbulent dynamo effects other than the $\\alpha$ effect.  { Numerical evidence for turbulent dynamo effects has so far mainly been obtained from convection simulations in small (compared to the dimensions of a star) rectangular boxes \\citep[e.g., ][]{braetal90,ossstibra01,ossetal02,giezierud05,kapetal06,cathug06,hugcat08}. Due to the assumption of a uniform mean magnetic field and other limitations, most of these studies could only find parts of the MEMF that are proportional to the LSMF, i.e., the $\\alpha$ effect and turbulent pumping (a contribution to the MEMF of the form $\\vec{\\mathcal{E}}_\\gamma=\\vec{\\gamma}\\times\\langle\\vec{B}\\rangle$, with $\\vec{\\gamma}$ denoting a vector; it leads to an advection of the mean magnetic field). Recently, however, \\citet{kapkorbra09}, who used a procedure referred to as the test field method \\citep{schrietal05,schrietal07} together with numerical simulations of turbulent convection with shear and rotation, were able to also identify the action of the combined $\\vec{\\Omega}\\times\\vec{J}$ and shear-current effects. }  Here, we explore axisymmetric kinematic dynamo models containing the $\\vec{\\Omega}\\times\\vec{J}$ and shear-current effects, differential rotation, and meridional circulation. In calculating the MEMF we use analytical expressions derived by \\cite{pip08} on the basis of a simplified version of the $\\tau$ approximation {\\citep[cf.][]{vaikic83,brasub05,brasub05b}}. We construct models with distributed dynamo action in the bulk of the convection zone, rather than in the overshoot layer at the bottom of the convection zone. The model calculations are carried out using the rotation profile of the Sun as obtained from helioseismic measurements and radial profiles of other quantities according to a standard model of the solar interior.  The remainder of the paper is organized as follows: In Sect.~\\ref{sec_model} we describe our dynamo model, as well as the used numerical procedure (some benchmark tests for our computer code are presented in Appendix \\ref{appendix_benchmarks}). Then, in Sect.~\\ref{sec_results}, we present the obtained results. In Sect.~\\ref{sec_conclusions}, we draw conclusions and discuss our results. ",
        "conclusions": "\\label{sec_conclusions}  We have studied kinematic axisymmetric mean-field dynamo models  in the geometry of a spherical shell, as appropriate for the Sun and solar-type stars, where the $\\vec{\\Omega}\\times\\vec{J}$ and shear-current effects were included as turbulent sources of the large-scale magnetic field while the $\\alpha$ effect was omitted. Besides the turbulent dynamo mechanisms and differential rotation, a meridional circulation, in the form of two stationary circulation cells, one in the northern and one in the southern hemisphere, also was incorporated into the models. We have concentrated on the dynamo onset.  Our results show that the $\\vec{\\Omega}\\times\\vec{J}$ and shear-current effects can, at least in principle, take over the role that the $\\alpha$ effect usually plays in mean-field dynamo models. However, only if the meridional flow is sufficiently fast are the characteristic properties of solar-type dynamos qualitatively correctly reproduced. In particular, the amplitude of the meridional flow needs to exceed a threshold value in order that the most unstable magnetic mode has dipolar parity and oscillates. This mode then also shows a latitudinal drift towards the equator within each half cycle and a phase relation between the poloidal and toroidal parts of the field in accordance with the observations of solar activity. The threshold value for the amplitude of the meridional flow, $u_0\\gtrsim 10\\,\\mathrm{m/s}$ (reached at the surface, the flow speed at the bottom is on the order of $1\\,\\mathrm{m/s}$), is consistent with solar observations and agrees with the value of $10\\,\\mathrm{m/s}$ often adopted in studies of flux-transport dynamos with the $\\alpha$ effect \\citep[cf., e.g.,][]{bonetal02,jouetal08}.  { In models of advection-dominated dynamos, the specifics of the turbulent dynamo mechanism that generates the mean poloidal field are less important than they are in models without meridional flow (given the rotation law and the generation of the mean toroidal field from the mean poloidal field by velocity shear). Once the field is generated, it is advected equatorwards by the flow. However, the distribution of the turbulent dynamo sources, or their more or less strong localization, decisively influences the parity properties of the LSMF.} Studies of flux-transport dynamos with an $\\alpha$ effect as the turbulent source of the LSMF indicate that the $\\alpha$ effect must be strongly localized at the bottom of the convection zone to ensure the correct (dipolar) parity of the LSMF \\citep{dikgil01,bonetal02}. In our models, the turbulent dynamo effects are distributed over the bulk of the convection zone, though they are strongest near the bottom of the included domain. {We note that our  turbulent dynamo sources are not introduced arbitrarily but are calculated using a standard model of the solar interior together with rotation rates obtained from helioseismic measurements. }  As other advection-dominated dynamo models, the models presented here work only if the effective magnetic diffusivity is strongly reduced compared to the mixing-length estimates. At a radial distance of, say, $0.85\\,R_{\\odot}$, the turbulent magnetic diffusivity in our models is about $0.2\\,\\eta_0$ ($\\eta_0=1.8\\cdot 10^9\\,\\mathrm{m}^2/\\mathrm{s}$ is the maximum value of the turbulent magnetic diffusivity in the convection zone according to the mixing-length estimate). Together with our value of {0.025} for the parameter $C_\\eta$ (which regulates the turbulence level), this gives an effective magnetic diffusivity of about {$10^7\\,\\mathrm{m}^2/\\mathrm{s}$}, in agreement with the upper limit of $3\\cdot 10^7\\,\\mathrm{m}^2/\\mathrm{s}$ for the turbulent magnetic diffusivity in the bulk of the convection zone given by \\citet{dikgil06,dikgil07} for flux-transport dynamos with the $\\alpha$ effect.  The cycle periods that we obtain are at least {three times} as long as the observed period of the solar activity cycle. Also, the ratio between the toroidal and poloidal parts of the large-scale magnetic field is significantly smaller than supposed for the solar convection zone (apparently a common problem of all flux-transport models). Here one should keep in mind that requiring a perfect fit to the solar details, as far as these are known, would overstress the models. For instance, a solution that bifurcates at the dynamo onset will change quantitatively if it is traced away from the bifurcation point. {Thus, ultimately, self-consistent nonlinear models will be needed.}"
    },
    "0910/0910.4424_arXiv.txt": {
        "abstract": "Obtaining accurate measurements of the initial mass function (IMF) is often considered to be the key to understanding star formation, and a universal IMF is often assumed to imply a universal star formation process. Here, we illustrate that different modes of star formation can result in the same IMF, and that, in order to truly understand star formation, a deeper understanding of the primordial binary population is necessary. Detailed knowledge on the binary fraction, mass ratio distribution, and other binary parameters, as a function of mass, is a requirement for recovering the star formation process from stellar population measurements. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.1806_arXiv.txt": {
        "abstract": "% We explore the relationship between radio-loud (RL) and radio-quiet (RQ) quasars using a set of optical/UV/X-ray measures that are quite independent of radio measures. We find RL sources to show larger average FWHM H$\\beta$, weaker FeII emission, no soft X-ray excess and no CIV blueshift--all characteristics manifested by a large fraction of RQ quasars (that we call Population A). We find that log L$_{1.4Ghz}$ =31.6 ergs s$^{-1}$ Hz$^{-1}$ (or R=70) is the lower limit for RL quasars showing FRII morphology. We find no evidence for a hidden FRII population below this level. We conclude that RL sources are a distinct quasar population that may also include 30-40\\% of RQ sources which apparently show similar geometry and kinematics (what we call Population B). This RQ overlap, if not coincidental, may include inactive RL quasars as well as quasars with geometry/kinematics similar to RL sources but where RL activity is inhibited in some way (e.g. host morphology, BH spin). ",
        "introduction": "Quasars were discovered because of their excess radio continuum emission yet, ironically, today we find only 7-10\\% of them to be radio-loud (RL). The vast majority ($\\sim$92\\%) of AGN are radio-quiet (RQ) One can see from the broad-band spectra displayed in NED that RL sources are dominated by a non-thermal power-law at all wavelengths. What is the relationship between the RL and RQ quasars? Are they indistinguishable at other wavelengths? Both RQ and RL quasars, for example, show broad emission lines (unless they are Blazars). Are the measured properties of the lines (e.g. width, equivalent width, line shape and rest frame displacement) similar?  Such questions can be difficult to answer using the majority of available spectroscopic data because at low resolution and s/n all quasars look alike. The advent of the Sloan Digital Sky Survey offers us more uniform, high resolution and high s/n data. This is  especially true if one restricts oneself to the brightest (g$\\leq$17) SDSS quasars.    In order to test the existence of a RL-RQ dichotomy we require an observational context within which to search for differences. This context should involve observational parameters that are  independent of radio measures. We have been developing a four dimensional parameter space that builds upon pioneering studies of the PG quasars (Boroson \\& Green 1992; Wang et al. 1996). This attempt at a multiwavelength  unification of quasar properties (4DE1) involves optical, UV and X-ray measures (Sulentic et al. 2000, 2007; Marziani et al. 2001, 2003a). Our four principal parameters are: 1) full width at half maximum of the broad H$\\beta$ emission line (FWHM H$\\beta$), 2) the equivalent width ratio of FeII4570 blend$\\over$H$\\beta$ (RFE), 3) the FWHM  normalized centroid shift of CIV1549 (C(1/2)) and 4) the soft X-ray photon index ($\\Gamma_{soft}$)  (Marziani et al. 2001, 2003a). We have addressed the RL-RQ question using both our own atlas sample (Marziani et al. 2003b) and a bright SDSS DR5 sample. Both are low z samples largely brighter than g=17.5 (Zamfir et al. 2008).     \\begin{figure} \\includegraphics[scale=1.04, bb = -20 0 250 215]{sulenticf1.eps} \\caption{Optical plane of the 4DE1 parameter space showing the distribution of the SDSS quasars FWHM of H$\\beta$ versus the FeII-prominence parameter RFE. Filled squares identify FRII sources, while the grey spot represent the general quasar population.} \\end{figure} ",
        "conclusions": ""
    },
    "0910/0910.3668_arXiv.txt": {
        "abstract": "We study the impact of theoretical uncertainty in the dark matter halo mass function and halo bias on dark energy constraints from imminent galaxy cluster surveys.  We find that for an optical cluster survey like the Dark Energy Survey, the accuracy required on the predicted halo mass function to make it an insignificant source of error on dark energy parameters is $\\approx 1\\%$.  The analogous requirement on the predicted halo bias is less stringent ($\\approx 5\\%$), particularly if the observable--mass distribution can be well constrained by other means. These requirements depend upon survey area but are relatively insensitive to survey depth.  The most stringent requirements are likely to come from a survey over a significant fraction of the sky that aims to observe clusters down to relatively low mass, $M_{\\rm th} \\approx 10^{13.7}\\ \\hiMsun$; for such a survey, the mass function and halo bias must be predicted to accuracies of $\\approx 0.5\\%$ and $\\approx 1\\%$, respectively.  These accuracies represent a limit on the practical need to calibrate ever more accurate halo mass and bias functions.  We find that improving predictions for the mass function in the low-redshift and low-mass regimes is the most effective way to improve dark energy constraints. ",
        "introduction": "Observations of the number density of galaxy clusters as a function of cluster mass and redshift are powerful probes of cosmology, especially the accelerated cosmological expansion caused by the cryptically dubbed {\\em dark energy}. Cosmological parameters have been recently inferred from a number of observations employing different techniques for cluster identification \\citep[e.g.,][]{Gladders07,Mantz08,Henry09,Rozo09,Vikhlinin09,Mantz09a}. The effort to derive precise and unbiased constraints on cosmological parameters, particularly those describing dark energy, requires control of various systematic error sources.  In this paper, we estimate the precision with which the abundance of dark matter halos as a function of mass (the halo {\\em mass function}) and the clustering of halos as a function of mass (the halo {\\em bias}) must be predicted in order to ensure that errors in the predictions for these quantities will be a small fraction of the error budget in forthcoming cluster count cosmology efforts.   For the current generation of cluster count surveys, such as the Sloan Digital Sky Survey (SDSS), the Red-Sequence Cluster Survey (RCS), and the Massive Cluster Survey (MACS), uncertainty in the mass function is unlikely a major concern.  The dominant errors in contemporary surveys are thought to be either limited statistics \\citep[e.g.,][for MACS]{Mantz08} or the uncertain relation between observable quantities and cluster mass \\citep[e.g.,][for RCS and SDSS respectively]{Gladders07,Rozo09}.  Uncertainty in the relation between cluster observables and masses will be an important limitation \\citep[e.g.,][]{Majumdar03,Majumdar04,LimaHu04,LimaHu05,Stanek06,LimaHu07,Cunha08}, and it is thought that some combination of empirical and theoretical insight into the observable--mass relation will continue to prove effective in the interpretation of forthcoming data \\citep[e.g.,][]{Kravtsov06, Rozo09Richness}.  Understanding the relation between observables and mass is a rapidly evolving subject, so it is difficult to anticipate the level at which this will be controlled in the analyses of forthcoming data.  Although we do not directly deal with observable--mass relations, we explore the dependence of our results on assumptions about the observable--mass relation.   The motivations to study the uncertainty in predicting halo abundances and clustering are twofold.  First, forthcoming optical surveys, such as the Dark Energy Survey (DES)\\footnote{ {\\tt http://www.darkenergsurvey.org}}, the Panoramic Survey Telescope \\& Rapid Response System (PanSTARRS)\\footnote{ {\\tt http://pan-starrs.ifa.hawaii.edu}}, and the Large Synoptic Survey Telescope (LSST)\\footnote{ {\\tt http://www.lsst.org}}; X-ray surveys, such as the extended ROentgen Survey with an Imaging Telescope Array (eRosita)\\footnote{ {\\tt http://www.mpe.mpg.de/projects.html\\#erosita}}; and Sunyaev--Zeldovich (SZ) surveys, such as the Atacama Cosmology Telescope (ACT)\\footnote{ {\\tt http://www.physics.princeton.edu/act/}} and the South Pole Telescope (SPT)\\footnote{ {\\tt http://pole.uchicago.edu/}}, all promise to expand greatly upon contemporary observations of clusters. These surveys will enable significant reductions in statistical errors and provide better constraints on the cluster observable--mass distribution.  With ever-improving theoretical understanding of this distribution, systematic errors that are currently unimportant may be damaging to future surveys and need to be controlled.  For example, \\citet{Crocce09} demonstrated that current errors in predicted mass functions may lead to statistically significant systematic errors in the inferred dark energy equation of state from SZ surveys.   Second, there is a significant, ongoing effort to develop ever more accurate predictions for the halo mass function \\citep[e.g.,][]{ShethTormen99,Sheth01,Jenkins01,Evrard02,Reed03, Warren06,Lukic07,CohnWhite08,Tinker08,Robertson09,Lukic09,Crocce09} and halo bias \\citep[e.g.,][]{ShethTormen99,Sheth01,SeljakWarren04,Tinker08,Tinker09}. Mass functions and halo biases have been derived on theoretical grounds \\citep[][see \\citealt{Zentner07} for a recent review] {PressSchechter74,Bond91,MoWhite96,ShethTormen99,Sheth01, Maggiore09a,Maggiore09b,Maggiore09c}.  Fits to numerical simulations with theoretically motivated functional forms have been provided by several groups, and the current state-of-the-art includes the recent papers by \\citet{Lukic07}, \\citet{Tinker08}, and \\citet{Crocce09}, as well as the forthcoming work of the Los Alamos National Laboratory group (S. Bhattacharya et al. 2010, in preparation).  Current theoretical predictions of the mass function are accurate at the $\\sim 10\\%-30\\%$ level \\citep[e.g.,][]{Tinker08,Robertson09,Rudd08,Stanek08}, and future efforts will continue to improve this accuracy. These efforts require numerous cosmological numerical simulations along with requisite verification, validation, and considerable analysis, involving large commitments of both computational and human resources.  Therefore, establishing the accuracy in the mass function and halo bias that are required by surveys is a particularly timely issue.   In this work, we quantify how accurately halo mass functions and biases need to be predicted in order to have a negligible contribution to the error budgets of forthcoming cluster surveys. We begin in Section~\\ref{sec:error} with a simple demonstration of the potential errors that may be induced by the inaccuracy in the mass function.  We present our mass function parameterizations and cosmological parameter constraint forecasts in Section~\\ref{sec:parameterization} and Section~\\ref{sec:implementation}, respectively.  In all cases, we take the \\citet{Tinker08} mass function and \\citet{Tinker09} halo bias as a fiducial model, about which we perturb to estimate errors relevant to forthcoming optical and SZ surveys.   We detail our results in Section~\\ref{sec:results}. In summary, we find that for relatively near-term optical surveys like DES, or SZ surveys like SPT, the mass function must be calibrated to the $\\approx 1\\%$ level while the halo bias must be calibrated to the $\\approx 5\\%$ level.  The requirement on the mass function is relatively insensitive to different assumptions about the observable--mass distributions within the range considered in contemporary work, whereas the requirement on the halo bias is more sensitive.  For longer-term experiments, such as a nearly half-sky optical survey planned for LSST, the calibration must be improved to the $\\approx 0.5\\%$ level in the mass function and the $\\approx 1\\%$ level in the halo bias.  This represents the most stringent requirement for the surveys being planned over the next decade and may serve as a practical endgame in the need to refine theoretical predictions for halo abundances and clustering.   We also explore the halo masses and redshifts at which it is most important to make accurate predictions for the halo mass function.  We find that the most effective strategy to improve dark energy constraints is to improve predictions at the low masses and low redshifts involved in the survey.  Other details of the dependence of calibration requirements on mass and redshift depend upon the observable--mass distribution.  We summarize our results and draw our conclusions in Section~\\ref{sec:conclusion}.   As we were completing our study, we learned of the related work of \\citet{Cunha09MF}. Although our results generally agree with the results in this study, our study has a number of distinct and complementary aspects. These authors concentrated on the specific parameters in the \\citet{Tinker08} mass function and the \\citet{ShethTormen99} halo bias, focusing considerable study on the degeneracies between nuisance parameters and cosmological parameters. In contrast, we aim to describe the uncertainty in mass function and halo bias in a manner that is independent of the form of any specific fitting function. We discretize the mass function and halo bias, and assign to each mass and redshift bin a distinct nuisance parameter specifying the fractional deviation of the mass function or halo bias from the fiducial values. This parameterization also allows us to identify the masses and redshifts at which it is most fruitful to refine predictions in order to limit systematic errors on dark energy parameters. ",
        "conclusions": "\\label{sec:conclusion}  We have studied the impact of theoretical uncertainties in predictions of the halo mass function and halo bias on dark energy constraints from upcoming cluster surveys. Our primary conclusions are as follows.   \\begin{enumerate}  \\item Inaccuracy in the shape of the mass function at the $20\\%$ level (similar to the current state of theoretical uncertainty) can lead to significant systematic errors in dark energy parameter inferences.  \\item For near-term cluster surveys like DES, the mass function must be predicted at approximately 1\\% accuracy to avoid more than 10\\% degradation in the resulting dark energy FoM. Similarly, the halo bias should be predicted with a precision of approximately 5\\%.  The current state of uncertainty in these functions could decrease the FoM of a DES-like survey by a factor of $\\sim 2$.  Requirements are generally less restrictive for SZ and X-ray efforts.  \\item A future optical survey over a significant fraction of the sky, like LSST, will require the most stringent constraints on the theoretical predictions.  In this case, the mass function and halo bias must be computed with an accuracy of $\\sim 0.5\\%$ and $\\sim 1\\%$ respectively in order to guarantee that the theoretical uncertainty in these quantities is a negligible contributor to the dark energy error budget.  This represents a practical limit to the accuracy with which these quantities will need to be computed in order to interpret future survey data.  \\item Precise prediction of the mass function at the low masses close to the observable threshold and low redshifts that will be used in the analysis of survey data is the most beneficial to improve dark energy constraints.  \\end{enumerate}   We have considered the influence of theoretical uncertainties in comparison only to statistical errors on forthcoming measurements; however, additional systematic errors that will be present at different levels in different types of cluster surveys will make the demands on theoretical mass functions somewhat less restrictive.  For example, the systematic errors due to inaccurate modeling of the observable--mass distribution \\citep[e.g. non-Gaussian scatter,][]{Cohn07, Shaw09}, an inaccurate calibration of completeness or false detection rate, and large errors in photometric redshift estimates will all increase errors on dark energy parameters and reduce the relative influence of uncertainties in predicted halo mass functions and biases.  The requirements we advocate are relevant when these other known sources of error do not dominate the error budget. A robust interpretation of our calculations is that the requirements we quote render the error due to inaccurate predictions of halo abundances unimportant, but less stringent requirements may be adequate when systematic errors are large.  However, we did show that the theoretical requirement is relatively insensitive to knowledge of the observable--mass relation when this relation is well described by a Gaussian distribution.  This suggests that the requirements may not be considerably loosened if systematics are moderate.   Dark matter simulations currently have $\\sim 5$\\% uncertainties in predicting mass functions for fairly standard cosmological models, even when the dark energy is {\\em fixed} to a cosmological constant \\citep{Tinker08}.  To achieve sub-percent level predictions in dark matter simulations requires a more careful consideration of systematics in halo finding \\citep{Heitmann05} and initial conditions \\citep{Crocce06}, as well as a more careful study of the non-universality of the halo mass function \\citep{Tinker08,Robertson09}.  This will require simulations of a wider range of cosmological models that go beyond the standard cold dark matter with cosmological constant model.  Reasonably exhaustive simulation programs are beginning \\citep[e.g.][]{Desjacques09,Maggiore09c,LamSheth09,Martino09,Pillepich08,Casarini09,Jennings09, Grossi09, Alimi09}. Even more challenging still will be understanding the systematic uncertainties induced by baryonic physics.  Recent studies have shown that these effects can result in significant deviations in the halo number density and that these effects are a function of the halo mass scale \\citep{Rudd08,Stanek08}.  Although meeting the stringent constraints we advocate may be quite challenging, pursuing further accuracy in theoretical predictions along with controlling various systematics in clusters surveys promises to continue to improve our knowledge of dark energy."
    },
    "0910/0910.1347_arXiv.txt": {
        "abstract": "We present a new method to retrieve molecular abundances and temperature profiles from exoplanet atmosphere photometry and spectroscopy. We run millions of 1D atmosphere models in order to cover the large range of allowed parameter space.  In order to run such a large number of models, we have developed a parametric pressure-temperature ($P$-$T$) profile coupled with line-by-line radiative transfer, hydrostatic equilibrium, and energy balance, along with prescriptions for non-equilibrium molecular composition and energy redistribution. The major difference from traditional 1D radiative transfer models is the parametric $P$-$T$ profile, which essentially means adopting energy balance only at the top of the atmosphere and not in each layer. We see the parametric $P$-$T$ model as a parallel approach to the traditional exoplanet atmosphere models that rely on several free parameters to encompass unknown absorbers and energy redistribution. The parametric $P$-$T$ profile captures the basic physical features of temperature structures in planetary atmospheres (including temperature inversions), and fits a wide range of published $P$-$T$ profiles, including those of solar system planets. We apply our temperature and abundance retrieval method to the atmospheres of two transiting exoplanets, HD~189733b and HD~209458b, which have the best {\\it Spitzer} and {\\it HST} data available. For HD~189733b, we find efficient day-night redistribution of energy in the atmosphere, and molecular abundance constraints confirming the presence of H$_2$O, CO, CH$_4$ and CO$_2$. For HD~209458b, we confirm and constrain the day-side thermal inversion in an average 1D temperature profile. We also report independent detections of H$_2$O, CO, CH$_4$ and CO$_2$ on the dayside of HD~209458b, based on six-channel {\\it Spitzer} photometry. We report constraints for HD~189733b due to individual data sets separately; a few key observations are variable in different data sets at similar wavelengths. Moreover, a noticeably strong CO$_2$ absorption in one data set is significantly weaker in another. We must, therefore, acknowledge the strong possibility that the atmosphere is variable, both in its energy redistribution state and in the chemical abundances. ",
        "introduction": "Major advances in the detection and characterization of extrasolar planet atmospheres have been made during the last decade. Over a dozen hot Jupiter atmospheres have been observed by the {\\it Spitzer Space Telescope} and a handful by the {\\it Hubble Space Telescope} ({\\it HST}). These observations have changed the field of exoplanets, enabling for the first time studies of atmospheres of distant worlds.  Observational highlights include the identification of molecules and atoms, and signatures of thermal inversions. {\\it HST} observed water vapor and methane in transmission in HD~189733b (Swain et al. 2008b). Water vapor was also inferred in transmission at the 3.6 $\\micron$, 5.8 $\\micron$, and 8 $\\micron$ in HD~189733b (Tinetti et al. 2007) (but c.f. Desert et al. 2009, and also Beaulieu et al. 2009). Water was inferred on the dayside of HD~189733b from the planet-star flux contrast in the 3.6 $\\micron$ - 8.0 $\\micron$ {\\it Spitzer} broadband photometry (Barman 2008; Charbonneau et al. 2008). Additionally, Grillmair et al (2008) reported a high signal-to-noise spectrum of the dayside of HD~189733b in the 5 $\\micron$ - 14 $\\micron$ range, using the {\\it Spitzer} Infrared Spectrograph (IRS), and reported detection of water. More recently, {\\it HST} detected carbon dioxide in thermal emission in HD~189733b (Swain et al. 2009a), in 1.4 $\\micron$ - 2.6 $\\micron$. Previously, sodium was detected in HD~209458b (Charbonneau et al. 2002) and HD~189733b (Redfield et al. 2008). Several observations of the planetary dayside have also  been reported for HD~209458b, first observed by {\\it Spitzer} at 24 $\\micron$ (Deming et al. 2005 and Seager et al. 2005). A major observational discovery was the finding of strong emission features in the broadband photometry of HD~209458b at secondary eclipse (Knutson et al. 2008), implying a thermal inversion in the atmosphere (Burrows et al. 2007 \\& 2008). In tandem with our present work, water, methane and carbon dioxide were independently inferred and recently reported on the dayside of HD~209458b, with {\\it HST} NICMOS and five-channel {\\it Spitzer} photometry (Swain et al. 2009b). Additionally, there are hints that variability of hot Jupiter thermal emission may be common (Knutson et al. 2007 vs. Charbonneau et al. 2008; Charbonneau et al. 2008 vs. Grillmair et al. 2008; Deming et al. 2005 vs. Deming, private communication 2009). The discoveries of atmospheric constituents and temperature structures mark the remarkable successes of exoplanet atmosphere models in interpreting the observations.  Despite the successes of model interpretation of data, major limitations of traditional self-consistent atmosphere models are becoming more and more apparent. Model successes include the detection of thermal inversions on the day-side of hot Jupiters (Burrows et al. 2007; Knutson et al. 2008), and subsequent classification of systems on the same basis (Burrows et al. 2008; Fortney et al. 2008; Seager et al. 2008). The nature of the absorbers causing the inversions, however, is not yet known (Knutson et al. 2008 \\& 2009a), forcing modelers to use an adhoc opacity source to explain the data. The intially favored opacity sources, TiO and VO may be unlikely to cause inversions of the observed magnitude (Spiegel et al. 2009).  A second model limitation arises because 1D radiative-convective equilibrium models cannot include the complex physics involved in hydrodynamic flows. Existing models use a parameterization of energy transfer from the day side to night side (e.g., Burrows et al. 2008), using parameters for the locations of energy sources and sinks, and for the amount of energy transported.  One further example of atmospheric model limitations involves the treatment of chemical compositions. Self-consistent models generally calculate molecular mixing ratios based on chemical equilibrium along with the assumption of solar abundances, or variants thereof (e.g., Seager et al. 2005). However, hot Jupiter atmospheres should host manifestly non-equilibrium chemistry (e.g., Liang et al. 2003, Cooper \\& Showman, 2006), which render the assumption of chemical equilibrium only a fiducial starting point.  With parameters used to cover complicated or unknown physics or chemistry, existing ``self-consistent'' radiative-convective equilibrium models are no longer fully self consistent. Furthermore, given the computational time of such models, only a few models are typically published to interpret the data, often only barely fitting the data ---a quantitative measure of fit is typically absent in the literature.  In an ideal world, one would construct fully self-consistent 3D atmosphere models (e.g., Showman et al. 2009) that include hydrodynamic flow, radiative transfer, cloud microphysics, photochemistry and non-equilibrium chemistry, and run such models over all of parameter space anticipated to be valid. Such models will remain idealizations until computer power improves tremendously, given that atmospheric circulation models can take weeks to months to run even with simple radiative transfer schemes.  We are motivated to develop a data-interpretation framework for exoplanet atmospheres that enables us to run millions of models in order to constrain the full range of pressure-temperature ($P$-$T$) profiles and abundance ranges allowed by a given data set. In this work, we report a new approach to 1D modeling of exoplanet atmospheres. Taking a cue from parameterized physics already being adopted by existing models, we go much further by parameterizing the $P$-$T$ profile as guided by basic physical considerations and by properties of $P$-$T$ profiles of planetary atmospheres in the literature. Indeed, it appears at present, the only way to be able to run enough models to constrain the $P$-$T$ structure and abundances is to use a parameterized $P$-$T$ profile. The essential difference between our new method and currently accepted radiative-convective equilibrium models is in the treatment of energy balance. We ensure global energy balance at the top of the atmosphere, whereas conventional atmosphere models use layer-by-layer energy balance by way of radiative and convective equilibrium. Our new method can be used as a stand-alone model, or it can be used to identify the parameter space in which to run a reasonable number of traditional model atmospheres. We call our method ``temperature and abundance retrieval\" of exoplanet atmospheres.  Atmospheric temperature and abundance retrieval is not new (e.g. Goody \\& Yung, 1989). Studies of planetary atmospheres in the solar system use temperature retrieval methods, but in the present context there is one major difference. Exquisite data for solar system planet atmospheres means a fiducial pressure - temperature profile can be derived. The temperature retrieval process for atmospheres of solar system planets therefore involves perturbing the fiducial temperature profile. For exoplanets, where data are inadequate, there is no starting point to derive a fiducial model. Nevertheless, Swain et al. (2008b, 2009a) used previously published model $P$-$T$ profiles of HD~189733b, and varied molecular abundances to report model fits to {\\it HST} NICMOS spectra of HD~189733b. In another study, Sing et al. (2008) used a parametric temperature profile with a thermal inversion, and parametrized abundances, required specifically to explain {\\it HST} STIS optical observations of HD~209458b. Our model, however, is completely general in the choice of $P$-$T$ profiles, ranges over tens of thousands of profiles, and satisfies the physical constraints of global energy balance and hydrostatic equilibrium. Our model reports, with contour plots, the quantitatively allowed ranges of $P$-$T$ profiles and molecular abundances. In addition, we also report constraints on the albedo and day-night energy redistribution, and on the effective temperature.  In this paper, the $P$-$T$ profiles are introduced in \\S~2, and the radiative transfer model and parameter space search strategy are described in \\S~3. Results of data interpretation with the model are presented in \\S~4 for HD~189733b, and in \\S~5 for HD~209458b. \\S~6 discusses the overall approach, and prospects for future work. And, in \\S~7 we present a summary of our results and conclusions. ",
        "conclusions": "\\label{sec-summary}  We have presented a solution to the retrieval problem in exoplanet atmospheres: given a set of observations, what are the constraints on the atmospheric properties?  Our method consists of a new model involving line-by-line radiative transfer with a parametric pressure-temperaure ($P$-$T$) profile, along with physical constraints of hydrostatic equilibrium and global energy balance, and allowing for non-equilibrium molecular compositions. This approach enabled us to run millions of models of atmospheric spectra over a grid in the parameter space, and to calculate goodness-of-fit contours for a given data set. Thus, given a set of observations and a desired measure of fit, our method places constraints on the $P$-$T$ profile and the molecular abundances in the atmosphere.  We used our method to place constraints on the atmospheric properties of two hot Jupiters with the best available data: HD~189733b and HD~209458b. We computed $\\sim 10^7$ models for each system. Our constraints on the $P$-$T$ profiles confirm the presence of a thermal inversion on the dayside of HD~209458b, and the absence of one in HD~189733b, as has been previously reported in the literature. And, the constraints on molecular abundances confirm the presence of H$_2$O, CH$_4$, CO and CO$_2$ in HD~189733b, as has been found by several other studies in the past.  We presented some new findings. We report independent detections of H$_2$O, CH$_4$, CO$_2$, and CO, from six-channel {\\it Spitzer} photometry of HD~209458b, at the $\\xi^2 < 1$ level (also see Swain et al. 2009b). At the $\\xi^2 = 2$ surface, we find the lower-limits on the mixing ratios of CH$_4$, CO and CO$_2$ to be $10^{-8}$, $4 \\times 10^{-5}$ and $2 \\times 10^{-9}$, respectively. In addition, we find an upper-limit on H$_2$O mixing ratio of $10^{-4}$. Secondly, our results using the high S/N IRS spectrum of HD 189733b indicate that the best fit models are consistent with efficient day-night circulation on this planet, consistent with the findings of Knutson et al. (2007 \\& 2009b); although, we also find that the broad band photometry data requires much less efficient day-night circulation, as has been reported by several studies in the past. The different constraints placed by the different data sets arise from the variability in the data.  A few key observations of HD~189733b vary between different data sets at similar wavelengths. Moreover, a noticeably strong CO$_2$ absorption in one data set is significantly weaker in another. We must, therefore, acknowledge the strong possibility that the atmosphere is variable, both in its energy redistribution state and in the chemical abundances.  Our new temperature and abundance retrieval --- together with the limited but remarkable data sets --- charts a new course for exoplanet atmosphere interpretation. Whether our approach is used as a stand alone model, or as a starting point for forward models, the field of exoplanet atmospheres is ready for a technique to place detailed atmospheric constraints governed by the data. The ideal data set is one with a wide wavelength range that covers more than one spectral feature of different molecules. The {\\it James Webb Space Telescope}, scheduled for launch in 2013, is ideal for observations of transiting exoplanet atmospheres and should usher a new era for quantitative analyses of exoplanet atmospheres."
    },
    "0910/0910.3174_arXiv.txt": {
        "abstract": "A nearby core collapse supernova will produce a burst of neutrinos in several detectors worldwide.  With reasonably high probability, the Earth will shadow the neutrino flux in one or more detectors. In such a case, for allowed oscillation parameter scenarios, the observed neutrino energy spectrum will bear the signature of oscillations in Earth matter.  Because the frequency of the oscillations in energy depends on the pathlength traveled by the neutrinos in the Earth, an observed spectrum contains also information about the direction to the supernova. We explore here the possibility of constraining the supernova location using matter oscillation patterns observed in a detector.  Good energy resolution (typical of scintillator detectors), well known oscillation parameters, and optimistically large (but conceivable) statistics are required. Pointing by this method can be significantly improved using multiple detectors located around the globe.  Although it is not competitive with neutrino-electron elastic scattering-based pointing with water Cherenkov detectors, the technique could still be useful. ",
        "introduction": "The core collapse of a massive star leads to emission of a short, intense burst of neutrinos of all flavors.   The time scale is tens of seconds and the neutrino energies are in the range of a few tens of MeV.  Several detectors worldwide, both current and planned for the near future, are  sensitive to a core collapse burst within the Milky Way or slightly beyond~\\cite{Scholberg:2007nu}.  The first electromagnetic radiation is not expected to emerge from the star for hours, or perhaps even a few days.  Therefore any directional information that can be extracted from the neutrino signal will be advantageous to astronomers who can use such information to initiate a search for the visible supernova.  We note that not every core collapse may produce a bright supernova: some supernovae may be obscured, and some core collapses may produce no supernova at all, in which case directional information will aid the search for a remnant.  The possibility of using the neutrinos themselves to point back to the supernova has been explored in the literature~\\cite{Beacom:1998fj, Tomas:2003xn}.  Triangulation  based on relative timing of neutrino burst signals was also considered in~\\cite{Beacom:1998fj}; however available statistics, as well as considerable practical difficulties in prompt sharing of information, makes time triangulation more difficult. Leaving aside the possibility of a TeV neutrino signal~\\cite{Tomas:2003xn} (which would likely be delayed), the most promising way of using the neutrinos to point to a supernova is via neutrino-electron elastic scattering: neutrinos interacting with atomic electrons scatter their targets within a cone of about 25$^\\circ$ with respect to the supernova direction.  The quality of pointing goes as $\\sim N^{-1/2}$, where $N$ is the number of elastic scattering events. In water and scintillator detectors,  neutrino-electron elastic scattering represents only a few percent of the total signal, which is dominated by inverse beta decay $\\bar{\\nu}_e + p \\rightarrow n + e^+$, for which anisotropy is weak~\\cite{Vogel:1999zy}.  Furthermore the directional information in the elastic scattering signal is available only for water Cherenkov detectors, for which direction information is preserved via the Cherenkov cone of the scattered electrons. Taking into account the near-isotropic background of non-elastic scattering events~\\cite{Tomas:2003xn}, a Super-K-like detector~\\cite{Ikeda:2007sa} (22.5~kton fiducial volume) will have 68\\% (90\\%) C.L.  pointing of about 6$^\\circ$  (8$^\\circ$) for a 10~kpc supernova; this could improve to $< 1^\\circ$ for next-generation Mton-scale water detectors. Long string water detectors~\\cite{Halzen:1995ex} do not reconstruct supernova neutrinos event-by-event and so cannot use this channel for pointing. Scintillation light is nearly isotropic and so scintillation detectors have very poor directional capability, although there is potentially information in the relative positions of the inverse beta decay positron and neutron vertices~\\footnote{This technique was studied for gadolinium-loaded scintillator in references~\\cite{Apollonio:1999jg,Hochmuth:2007gv}; however gadolinium loading for future large scintillator detectors does not seem to be planned.}, and some novel scintillator directional techniques are under development~\\cite{nutech}.  We consider here a new possibility: detectors with sufficiently good energy resolution will be able to obtain directional information by observing the effects of neutrino oscillation on the energy spectrum of the observed neutrinos, assuming that oscillation parameters are such that matter oscillations are present.  Although not competitive with elastic scattering, some directional information can be obtained even in a single detector (unlike for time triangulation).  Combinations of detectors at different locations around the globe may yield fairly high quality information. ",
        "conclusions": "We have assumed in these idealized scenarios perfect knowledge of oscillation parameters. In practice, imperfect knowledge of the oscillation parameters will create some uncertainties.  In particular, the power spectrum peak position is sensitive to the value of $\\Delta m^2_{12}$; the peak height is sensitive to both $\\Delta m^2_{12}$ and $\\theta_{12}$; $\\theta_{13}$ also has an effect on both $k_{\\rm peak}$ and $h$. (The oscillation pattern is quite insensitive to the 23 mixing parameters.) Figs.~\\ref{fig:parameffect1} and \\ref{fig:parameffect2} show the effect on $k_{\\rm peak}$ and $h$ values of varying the oscillation parameters within currently allowed ranges~\\cite{Abe:2008ee}.  \\begin{figure}[!htbp]  \\includegraphics[width=3.2in]{otherparams_k.pdf}  \\caption{Effect of varying  $\\Delta m^2_{12}$ on the peak position as a function of $L$; in each case other oscillation parameters are held at their nominal values.} \\label{fig:parameffect1} \\end{figure}   \\begin{figure}[!htbp]  \\includegraphics[width=3.2in]{otherparams_h.pdf}  \\caption{Effect of varying $\\theta_{12}$ and $\\theta_{13}$ on the peak height as a function of $L$; in each case the other oscillation parameters are held at their nominal values. } \\label{fig:parameffect2} \\end{figure}   From these plots one can infer that $\\lsim$1\\% knowledge of the mixing parameters is desirable.  However, one can be quite optimistic that such precision will have been attained by the time a core collapse supernova happens when a large scintillator detector is running.  Another uncertainty that will affect the quality of pointing is that of the density of matter in the Earth.  We found only small differences in $k_{peak}$ and $h$ from varying the mantle density by $\\pm 3\\%$, or from varying the overall density by $\\pm 5\\%$, but observed some changes in peak pattern for the case of neutrinos passing through the core when varying the core density by $\\pm 10\\%$.   Many other effects may degrade the quality of direction information that can be obtained using this technique.  There may be real spectral features (\\textit{e.g.} ``splits'') which introduce additional Fourier components that could mask the peak, and detector imperfections may do the same.  We acknowledge also that there may be practical difficulties with the rapid exchange of information between experimenters required for prompt extraction of directional information from multiple detectors.  Nevertheless this technique represents an interesting possibility-- even half the sky is better than no directional information.  The oscillation pattern gives information about direction with even a single detector, and enhances any multiple-detector time-triangulation information.  Even if information from only a single detector is available, or if there are significant ambiguities, one can imagine also looking at the intersection of the allowed region with the Galactic plane regions for which supernovae are most likely to occur (\\cite{Mirizzi:2006xx}) (perhaps using the known probability distribution as a Bayesian prior) to improve the chances of finding the supernova: see Fig.~\\ref{fig:sky_coverage_vs_declination_with_mirizzi}.  \\begin{figure}[!htbp] \\begin{centering} \\includegraphics[width=3.2in]{1det_enhanced_skycov_declination_scint_mirizzi.pdf} \\caption{Average scintillator sky coverage vs declination for a single detector in Finland, with the expected probability for supernova occurrence superimposed from reference~\\cite{Mirizzi:2006xx}.} \\label{fig:sky_coverage_vs_declination_with_mirizzi} \\end{centering} \\end{figure}  We note that these estimates of pointing quality have been done using a fairly simple technique based on only two parameters characterizing the power spectra.  One can imagine employing more sophisticated algorithms, \\textit{e.g.} making use of secondary peaks or matching to a template, and possibly incorporating knowledge of specific detector properties or neutrino flux spectral features.  So although real conditions may degrade quality, with this simplified study we have not fully exploited all potentially available information.  As a final note: the technique could in principle work to determine directional information for neutrino signals from other astrophysical sources, such as black hole-neutron star mergers~\\cite{Caballero:2009ww}, assuming sufficient statistics."
    },
    "0910/0910.3945_arXiv.txt": {
        "abstract": "The evidence from velocity-resolved reverberation mapping showing a net infall of the broad-line region (BLR) of AGNs is reviewed.  Different lines in many objects at different epochs give a consistent picture of BLR motions.  The motions are dominated by virialized motion (rotation plus turbulence) with significant net inflow.  The BLR mass influx is sufficient to power the AGN.  The increasing blueshifting of lines with increasing ionization potential is a consequence of scattering off infalling material.  The high blueshiftings of the UV lines in Narrow-Line Seyfert 1s are due to enhanced BLR inflow rates rather than strong winds.  Seemingly conflicting cases of apparent outflow reverberation mapping signatures are a result of the breakdown of the axial-symmetry assumption in reverberation mapping.  There are several plausible causes of this breakdown:  high-energy variability tends to be intrinsically anisotropic, regions of variability are necessarily located off-axis, and  X-ray observations reveal major changes in line-of-sight column densities close to the black hole.  Results from reverberation mapping campaigns dominated by a single event need to be treated with caution. ",
        "introduction": "At a conference on ``Accretion and Ejection in AGNs'' one could na\\\"ively expect roughly equal time for both accretion and ejection!  However, inspection the AGN literature (and the conference program) shows that, while the theory of ejection in jets and winds, and observations of ejection from $\\gamma$-rays to radio waves are all well covered, there is almost nothing on observations of {\\em inflow}.  Yet we all believe that AGNs are ultimately powered by accretion onto black holes, so {\\em something} has to be inflowing!  I argue here that the ``something'' is the broad-line region (BLR) and that we can see consequences of the inflow. ",
        "conclusions": ""
    },
    "0910/0910.2458_arXiv.txt": {
        "abstract": "We present an initial survey of Mg {\\sc II} absorption characteristics in the halos of a carefully constructed, volume-limited subsample of galaxies embedded in the spectroscopic part of the Sloan Digital Sky Survey.  We observed quasars near sightlines to 20 low-redshift ($z \\sim 0.1$), luminous (\\Mrh~$\\leq -20.5$) galaxies in SDSS DR4 and DR6 with the LRIS-B spectrograph on the Keck I telescope.  The primary systematic criteria for the targeted galaxies are a redshift $z \\gtrsim 0.1$ and the presence of an appropriate bright background quasar within a projected 75 h$^{-1}$ kpc of its center, although we preferentially sample galaxies with lower impact parameters and slightly more star formation within this range.  Of the observed systems, six exhibit strong [$\\mew \\geq 0.3$~\\AA] Mg {\\sc II} absorption at the galaxy's redshift, six systems have upper limits which preclude strong \\MgII\\ absorption, while the remaining observations rule out very strong [$\\mew \\geq 1-2$~\\AA] absorption. The absorbers fall at higher impact parameters than many non-absorber sightlines, indicating a covering fraction $f_c \\lesssim 0.4$ for $\\geq 0.3$~\\AA\\ absorbers at $z \\sim 0.1$, even at impact parameters $\\leq$35 h$^{-1}$ kpc (f$_{\\rm c} \\sim 0.25$).  The data are consistent with a possible dependence of covering fraction and/or absorption halo size on the environment or star-forming properties of the central galaxy. ",
        "introduction": "\\label{sec:intro}  Because gas consumption timescales in galaxies are short, massive galaxies must have large reservoirs of gas in order to fuel star formation throughout cosmic time.  These reservoirs are likely quite diffuse and extended, and are therefore difficult to probe directly. Absorption lines in background quasars near galaxy sightlines allow a unique probe of this gas \\citep[e.g.,][and many more]{Bergeron91,Steidel92,Steidel94,Steidel95,Chen01}.  Absorption-line studies in transitions like \\MgII\\ have revealed a wealth of information about the absorbing gas and its relationship to galaxies \\citep[e.g.,][]{Bergeron91, Steidel94, Churchill96, Nestor05,Zibetti05, Nestor06,Nestor07,Zibetti07}.  These results suggest that many --- if not all --- luminous intermediate-redshift galaxies are surrounded by \\MgII-enriched gas that extends out to at least $\\sim$40 h$^{-1}$ kpc for luminous systems \\citep[$\\sim$ 42 h$^{-1}$ kpc if $\\Omega_{\\Lambda} = 0.7$;][]{Steidel95} and likely much further in some cases \\citep{Churchill05, Kacprzak08}.  Estimates of the covering fractions of these gaseous halos range from $0.17 < f_c \\leq 1$ \\citep{Steidel95, Bechtold92, Bowen95, Kacprzak08} in ``strong'' [\\ew\\ $> 0.3 \\AA$] absorption systems.  This threshold for strong absorbers is common and physically based; strong systems differ in character from weaker systems, which have higher covering fractions and different kinematics \\citep{Churchill99,Churchill05, Nestor05, Nestor06}.  If gaseous halos around galaxies are reservoirs for star formation \\citep{Maller04}, the observed properties of galaxies likely depend on whether their gas halos are intact.  Processes that may inhibit star formation in satellite galaxies include ``strangulation,'' where the outer hot gas halo of a satellite galaxy is removed by the surrounding environment \\citep[e.g.,][]{Larson80,Kawata08}. This removal prevents the further infall of cool gas.  A more violent and immediate process is ram pressure stripping, which acts to remove cold, star-forming gas from galaxies, possibly quenching star formation soon after infall \\citep[e.g.,][]{Gunn72,Moore96}.  These processes likely play a role in the observed environmental dependence of galaxy properties like color and morphology \\citep[e.g.,][]{Gunn72,Dressler80,Postman84, Blanton05a}. If stripping of the outer gas reservoir does regulate star formation, we might expect to find correlations between the properties of the galaxies and those of their halos.  Many \\MgII\\ studies find dependence on galaxy luminosity or halo mass \\citep[e.g.,][]{Steidel94, Bouche06}.  Recently, \\citet{Kacprzak07} have shown that stronger absorbers also preferentially reside in asymmetric galaxies.  Many early studies did not reveal trends with galaxy color or star-forming properties \\citep{Steidel94,Guillemin97}, although at least one study suggested that trends are possible \\citep{Bowen95}.  More recently, \\citet{Zibetti07} study a sample of nearly 3000 \\MgII\\ systems with \\ew\\ $> 0.8$~\\AA\\ at $0.37 < z < 1$ by using image-stacking techniques to probe the amount, spatial distribution, and colors of the integrated galaxian light from the multiple, stacked absorbers.  Their results demonstrate that \\MgII\\ absorption is color dependent: the integrated galaxian light is bluer for stronger absorbers and for galaxies that are closer to the quasar probes. However, these results are aggregate and are luminosity-weighted; it is difficult to disentangle the possible effects that result if, say, \\MgII\\ absorption is more common in lower-luminosity blue galaxies. Thus, while their creative statistical method has proven extremely powerful, it cannot reveal a one-to-one relationship that may exist between galaxy and \\MgII\\ properties.  \\begin{figure} \\plotone{figure1.eps} \\caption{Constructing a volume-limited sample of galaxies to probe for halo \\MgII\\ absorption. We illustrate the challenge inherent in the use of SDSS to construct a sample for ground-based observations of halo\\MgII\\ absorption. ({\\it Bottom}) We show the color-magnitude distribution of a volume-limited sample to -18.5, showing both red and blue galaxies; observing as faint as possible is desirable to sample the blue cloud.  As a function of absolute $r$-band magnitude limit chosen from the SDSS, we show the upper redshift limit ({\\it Middle}) and the expected \\MgII\\ absorption wavelength ({\\it Top}).  As we demonstrate in this study, wavelengths above $\\sim$3050 are largely accessible from the ground, suggesting that volume-limited samples from SDSS with limits in the range of $\\lesssim -21$ to $-20$ are possible to probe.  Here, we focus on the \\Mrh $\\lesssim -20.5$ sample.} \\label{fig:vol_lim} \\end{figure}  \\begin{figure} \\plotone{figure2.eps} \\caption{Distribution of impact parameters probed in the study ({\\it Dotted line}) compared to a uniform distribution from 5-75 h$^{-1}$ kpc.  The sample is deficient in galaxies with impact parameters above $\\sim$50h$^{-1}$ kpc (P$_{\\rm K-S} = 0.01$).} \\label{fig:impact_histo} \\end{figure}  Most of our knowledge of \\MgII\\ absorbers comes from redshifts $\\gtrsim 0.4$, where the \\MgII\\ doublet ($\\lambda$ 2796.35,2803.53~\\AA) is redshifted into the optical. Systems in lower-redshift galaxies are difficult to observe from the ground and, as a result, have not been studied as extensively.  To date, detections indicate that the absorber abundances and covering fractions decline somewhat.  Depending on their \\ew, the decline is not necessarily dramatic as redshifts approach zero \\citep{Bowen95,Nestor05, Nestor06}.  The vast majority of existing \\MgII\\ studies of individual absorber-galaxy associations begin with quasar spectra and then use photometric or spectroscopic redshifts to identify the nearby galaxies responsible for the \\MgII\\ absorption \\citep[e.g.,][and many others]{Bergeron91,Steidel94, Churchill99}.  Studies that begin with quasar spectra frequently focus on galaxies already known to have \\MgII\\ absorption.  Some also tabulate galaxies that clearly lack absorption, but \\citet{Churchill05} note that focused ``control'' studies are biased toward quasars near non-absorbers, or absorbers with patchier \\MgII-absorbing gas. As \\citet{Tripp05} point out, ``reversing'' the problem by identifying a controlled galaxy sample first, and then finding a subset with background quasars, alleviates significant selection effects.  \\citet{Bowen95} present a Hubble Space Telescope study of very nearby galaxies and \\citet{Tripp05} present preliminary results from a ``reverse'' study at $z\\sim 0.2$ using galaxies selected photometrically from the SDSS.  Here, we present a ``reverse'' study from a well-defined, volume-limited spectroscopic sample of $\\sim$L$^{\\star}$ galaxies from the Sloan Digital Sky Survey (SDSS).  Their basic properties, such as environment, star-formation, and metallicity, can be placed into the very well-defined context of the SDSS.  In \\S~2 we describe our sample selection and observational approach, and present the data.  \\S~3 contains an analysis of our results, and we conclude in \\S~4. ",
        "conclusions": "\\label{sec:conclusion}  We use ground-based spectroscopy with LRIS-B on the Keck I telescope to evaluate the \\MgII\\ absorption characteristics of a sample of $\\sim$L$^{\\star}$ galaxies selected as cleanly as possible from a volume-limited subsample to M$_{\\rm r} + 5 \\log{h} \\gtrsim -20.5$ of the SDSS DR6 Value Added Galaxy Catalog \\citep{Blanton03}.  We target 20 systems, achieving a range in the quality of spectral observations that reveals 6 ``strong'' [\\ew\\ $> 0.3 $~\\AA] absorbers and 6 non-absorbers.  We find the following:  \\begin{enumerate}  \\item Ground-based observations of \\MgII\\ absorption are possible --- with relatively low S/N ratio --- to redshifts as low as $z \\sim 0.1$ with modest amounts of observing time.  \\item The overall covering fraction, $f_c$, of gas capable of creating strong ($\\geq 0.3$~\\AA) \\MgII\\ absorption must be $< 1$ at $z \\sim 0.1$, broadly consistent with some previous studies.  Even inside the traditional absorption halo radius of $\\sim$35 h$^{-1}$ kpc, the covering fraction of strong absorption is $< 1$, with a naive estimate of f$_{\\rm c} \\lesssim 0.25$.  \\item The \\MgII\\ detection in the quasar with the largest impact parameter lies 62 h$^{-1}$ kpc from the center of the associated galaxy.  Thus, at least in rapidly star-forming $z \\sim 0.1$ galaxies, strong, but possibly unfilled, \\MgII\\ absorption gas halos can extend to at least this distance from a luminous galaxy.  \\item The data suggest, at the $\\lesssim$2$\\sigma$ significance level, that the \\MgII\\ absorption properties of the outer halos of galaxies may depend on the larger-scale environment of the absorbing galaxy. Absorbers appear to favor low-density environments, while non-absorbers show a possible preference for denser regions.  This result is consistent with ``strangulation'' models in which satellite galaxies are red because their outer gas halos are stripped when they become substructure.  \\item The relationships between galaxy color, impact parameter, and absorbing halo \\ew\\ are consistent with a model in which red galaxies with \\MgII\\ absorption are typically LINERs which are observed at smaller impact parameters.  Blue galaxies are possible \\MgII\\ absorbers to very large impact parameters (at least 62 h$^{-1}$ kpc); in fact, all of the bluest galaxies in this study are either strong absorbers or inconclusive.  These color trends are consistent with the results from statistical image-stacking studies \\citep{Zibetti07}, and with the broad expectations of theoretical models.  \\end{enumerate}"
    },
    "0910/0910.5755.txt": {
        "abstract": "We investigate spherically symmetric perfect-fluid spacetimes and discuss the existence and stability of a dividing shell separating expanding and collapsing regions. We perform a $3+1$ splitting and obtain gauge invariant conditions relating the intrinsic spatial curvature of the shells to the Misner-Sharp mass and to a function of the pressure that we introduce and that generalizes the Tolman-Oppenheimer-Volkoff equilibrium condition.\\textbf{ }We find that surfaces fulfilling those two conditions fit, locally, the requirements of a dividing shell and we argue that cosmological initial conditions should allow its global validity\\textbf{. }We analyze the particular cases of the Lema\\^{\\i}tre-Tolman-Bondi dust models with a cosmological constant as an example of a cold dark matter model with a cosmological constant ($\\Lambda$-CDM) and its generalization to contain a central perfect-fluid core. These models provide simple, but physically interesting illustrations of our results. ",
        "introduction": "%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% Models of structure formation generally assume that small local inhomogeneities grow due to the gravitational instability, so that the overdensities collapse and eventually form the \\textquotedbl{}bound\\textquotedbl{} structures we observe in the present universe. Underlying this viewpoint is the idea that the collapse of the overdensities departs from the general expansion of the universe. This approach often relies on the idea that a small overdensity can be approached as a closed patch in an otherwise spatially flat Friedmann universe, and it claims that Birkhoff's theorem justifies that, on the one hand, its evolution is independent from the outside universe, and, on the other hand, that the behavior of the outside Friedman universe is immune to the collapse of the closed patch (see e.g. \\citep{Peebles 1981,Padmanabhan 1993,Cattoen:2005dx}).%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% The collapse of overdensities has been extensively studied and most works have been focused on the study of the formation both of small structure (astrophysical objects) and of large-scale structure as the outcome of the growth of small perturbations in a cosmological context. The latter subject comprises the relativistic and Newtonian analysis of the evolution of the fluctuations (see e.g. \\citep{deSitter:1933,Hawking:1966qi,Mukhanov:1990me,Liddle:1993fq}) and the study of the subsequent amplification of the growing modes into the nonlinear regime resorting to numerical methods (see e.g. \\citep{Bernardeau:2001qr,Mota:2004pa,LeDelliou:2005ig,Maor:2006rh}). %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% % \\begin{comment} The underlying assumption has always been that the universe surrounding the collapsing overdensity - the so-called background universe - evolves in a way which is independent of the collapsing process. \\\\  \\end{comment} {}In the present work we consider spherically symmetric, inhomogeneous universes with pressure and study the question of whether there exists a dividing shell separating expanding and collapsing regions. %\\textsl{From the astrophysical viewpoint, it is justified to investigate %the formation of bound structures or the ultimate fate of finite overdensities, %assumed to be surrounded by a very rarefied atmosphere that can be %taken as empty. One thus describes the latter to be vacuum or, at %least, by an asymptotically flat spacetime. In the more general setting %of cosmology, the endeavour of understanding the large-scale structure %formation places us in a similar situation, albeit a generalized one. %We aim at explaining the ultimate formation of bound structures of %a larger scale than in the former astrophysical scale and so it makes %sense to assume that the collapsing excess of matter is surrounded %by a flat Friedmann universe, and hence by a conformally flat spacetime.} %\\textsl{However there is one question which has been left aside in %doing so. We don't have a clear explanation for the scales both spatial %and material/energetic that select the actual bound structures apart %from the Jean's criterion holding on the linear approximation. Seemingly %we don't understand }\\textsl{\\textcolor{red}{... unfold Ellis question}}\\textsl{. %We believe that the question addressed in the present paper has not, %to our knowledge, been properly clarified in the literature.} Our goal bears a connection to the general problem of assessing the influence of global physics into the local physics \\citep{Ellis:2001cq,Faraoni:2007es}. One aspect of this problem that has always attracted great interest is the endeavor to explain the local inertial phenomena in a Machian sense (see e.g. \\citep{Sciama:1953zz,Dicke:1961ma}) and, in fact, Brans-Dicke theory \\citep{Brans:1961sx,Barrow:1994nx,Mimoso:1994wn,Iguchi:2004yh} stems from this problem.  Another related aspect has been the study of the influence of cosmic expansion on local systems. Einstein and Straus \\citep{ES} were the first to study this problem by constructing a global solution that resulted from matching the spherically symmetric vacuum Schwarzschild solution to an expanding dust Friedmann-Lema\u00eetre-Robertson-Walker (FLRW) exterior across a hypersurface preserving the symmetry. Bonnor has made several investigations along this line (see e.g. \\citep{Bonnor1974}). In particular, he copresented an exact solution representing a local distribution of %ECD ( electrically counterpoised dust %) embedded in an expanding universe with zero spatial curvature \\citep{Bonnor 1996}, showing that the distribution participates in the expansion.% \\begin{comment} % \\footnote{Pressureless charged matter in equilibrium is dubbed electrically counterpoised dust (ECD), because its charge density is equal to its mass density (in relativistic units).% } In \\citep{Faraoni:2007es}exact solutions describing a relativistic central object embedded in a Friedmann universe are obtained. It is shown that the {}``all or nothing'' behavior recently discovered (i.e. weakly coupled systems are comoving while strongly coupled ones resist the cosmic expansion) is limited to the de-Sitter background. \\end{comment} {} Among the generalizations of this model are settings that keep the spherical symmetry but generalize the interior source fields by considering, for example, Vaidya (see \\citep{Fayos} and references therein) or Lema\u00eetre-Tolman-Bondi (LTB) spacetimes (see \\citep{Krasinski,Krazinski 2001,Krasinski:2003yp,Hellaby:2005ut,Mena:2004ck}% \\begin{comment} for references concerning the latter in connection with formation of voids \\end{comment} {}). On a different context, Herrera and co-workers \\citep{Herrera92,DiPrisco94,Herrera} have studied % \\begin{comment} small anisotropic perturbations around spherically symmetric homogeneous fluids in equilibrium. They conclude that this may lead to instabilities that result in \\end{comment} {} the \\textquotedbl{}cracking\\textquotedbl{} of compact objects in astrophysics using small anisotropic perturbations around spherically symmetric homogeneous fluids in equilibrium. The latter references are concerned with the existence of a shell where there is a change in the direction of the radial force acting on the particles of the shells. Whenever this happens one has a cracking situation, a concept introduced by Herrera in Ref. \\citep{Herrera92}. The approach of these works is somewhat complementary to ours because it is not the full evolution that is depicted there, but rather the effect on particles as a result from a departure from equilibrium.  In this work we use a different approach from all the works described above. On one hand, by making use of a single coordinate patch, we do not have to handle the matching problem. On the other hand, our approach is not perturbative. We adopt the formalism that has recently been developed in a remarkable series of papers by Lasky and Lun using generalized Painlev\u00e9-Gullstrand (GPG) coordinates \\citep{Adler et al 05,LaskyLun06b,LaskyLun06+}. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% We perform a $3+1$ splitting and obtain gauge invariant conditions relating not only the intrinsic spatial curvature of the shells to the Misner-Sharp mass% \\footnote{also referred to as ADM mass %, after its first users when considering the mass of the whole spatial hypersurface.% } \\citep{MisnerSharp} but also a function of the pressure that we introduce and that generalizes the Tolman-Oppenheimer-Volkoff (TOV) equilibrium condition.  In particular, we consider that the existence of a spherical shell separating an expanding outer region from an inner region collapsing to the center of symmetry, depends essentially on two conditions. The first condition establishes that there is no matter exchange across the shell% \\begin{comment} precise balance between two energy quantities that are the analogues of the total and potential energies at the dividing shell. (This amounts to the vanishing of the kinetic energy of the shell.) \\end{comment} {}. The second condition establishes that the generalized TOV equation is satisfied on that shell, and hence that this shell is in some sort of equilibrium. The difference with respect to the original problem where the TOV equation was introduced for the first time is twofold. Our result does not rely on the assumption of a static equilibrium of the spherical distribution of matter, and consequently does not assume that all the internal spherical perfect-fluid spherical shells are constrained to satisfy the TOV equation. In our case the generalized TOV equation is just satisfied at the dividing shell. % \\begin{comment} On the neighbouring shells it won't be satisfied, and these shells will either be collapsing or expanding since they are not in equilibrium. \\end{comment} {}Besides, the generalized TOV function depends on the spatial 3-curvature in a more general way than the original TOV equation. Furthermore, we shall characterize the dividing shell with kinematic quantities that provide a gauge invariant formulation of the problem.  In order to illustrate our results we will analyze some particular cases. The simplest example is provided by the well-known Lema\\^{\\i}tre-Tolman-Bondi dust models with a cosmological constant that can be seen as an example of a $\\Lambda$-CDM model. A preliminary presentation of this work can be found in \\citep{Delliou:2009dm}. As a second case we consider generalizations of the previous model to contain a central perfect-fluid core. These models provide simple, but physically interesting illustrations of our results.% \\begin{comment} A short summary of these main ideas can be found in \\citet{LeDMM09let}. \\end{comment} {}  % \\begin{comment} We also show how the absence of a difference between the pressures of neighbouring shells prevents the existence of a separating shell in the spatially homogeneous FLRW models. Conversely, we discuss how the classical reasoning that considers the collapse closed FLRW patch embedded in a flat Friedmann spacetime (see e.g \\citet{Peebles 1981,Padmanabhan 1993}) fits into our analysis. \\end{comment} {}  %\\textsl{TOV related :\\citep{Bekenstein 1971} who \\textquotedbl{}obtains %the metric corresponding to an arbitrary spherically symmetric distribution %of charged perfect-fluid without restricting consideration to the %static case, and shows that it joins onto the standard Reissner-Nordstr\u00f6m %exterior solution of Einstein's equations\\textquotedbl{} (which is,the %latter, static!);}   %\\textsl{\\citep{Eisenstaedt 1975} \\textquotedbl{}Exact solutions of %the Einstein field equations in the spherical case are derived and %applied for constructing a model representing a spherical mass immersed %in a cosmological universe. A comoving, spatially homogeneous system %of coordinates is utilized, and a convenient definition of the mass %of the system permits the complete and exact integration of the field %equations. First and second-region solutions are obtained and classified %according to the values of two parameters, one which appears to be %related to the structure of space-time and another which is a generalization %of the Schwarzschild mass for an immersed sphere. Physical interpretations %of the solutions are discussed.\\textquotedbl{}}   %\\textsl{\\citep{Boonserm:2006up} \\textquotedbl{} Develops \\textquotedbl{}solution %generating\\textquotedbl{} theorems for the TOV, whereby any given %solution can be \\textquotedbl{}deformed\\textquotedbl{} to a new solution.(...) %He conclude by presenting the interpretation of the TOV equation as %an integrability condition on the density and pressure profiles. \\textquotedbl{}}   %\\textsl{Collapse of structure surrounded by static spacetime-matching %conditions: \\citep{Sharif 2006}, who \\textquotedbl{}investigates %the effect of a positive cosmological constant on spherically symmetric %collapse with perfect-fluid, giving matching conditions between static %exterior and non-static interior spacetimes in the presence of a cosmological %constant.\\textquotedbl{}}   %\\textsl{\\citep{Bonnor1974} who \\textquotedbl{}shows that there exist %spherically symmetric inhomogeneous models which approach open R-W %ones as $t\\to\\infty$, but whose density on an initial space-like %surface is an arbitrary function of the radial variable.\\textquotedbl{}}   %\\citep{Krazinski 2001} generalises this result of Bonnor by showing %that any two spherically symmetric density profiles in Lema\\^{\\i tre}-Tolman %models specified on any two constant time slices can be joined by %the evolution of these models.   %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%   %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%   An outline of the paper is as follows: Section II The GPG-formalism of Lasky and Lun: $3+1$ splitting and gauge invariants kinematical quantities. Section III Existence of a shell separating contraction from expansion: general conditions. Section IV Particular examples: Section IV A $\\Lambda$-CDM model (LTB with a cosmological constant). Section IV B Perfect-fluid core in a $\\Lambda$-CDM model. Section V Discussion of our results.  We shall use units such that $8\\pi G=1=c$, and the following index convention: Greek indices $\\alpha,\\beta,...=1,2,3$ while latin indices $a,b,...=0,1,2,3$. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "In the present work we have considered spherically symmetric, inhomogeneous universes in order to ascertain under which conditions a dividing shell separating expanding and collapsing regions exists. This endeavor is important in relation with the present understanding of structure formation as the outcome of gravitational collapse of overdense patches within an overall expanding universe% \\begin{comment} , since there is an underlying belief that the two disparate behaviours decouple, and is related to the assessment of the influence of global physics on local physics \\end{comment} {}.  We have addressed this problematic by resorting to an ADM 3+1 splitting, utilizing the so-called generalized Painlev\u00e9-Gullstrand coordinates as developed in Refs.~\\citep{Adler et al 05,LaskyLun06b}. This enables us to follow a nonperturbative approach and to avoid having to consider the matching of the two regions with the contrasting behaviors \\citep{Matravers:2000cu}. We have found local conditions characterizing the existence of a dividing shell. % \\begin{comment} One is a condition establishing the precise balance between two energy quantities that are the analogues of the total and potential energies at the dividing shell (this amounts to the vanishing of the kinetic energy of the shell). The second condition establishes that a generalized TOV equation is satisfied on that shell, and hence that this shell is in equilibrium. \\end{comment} {} We have related these conditions to a gauge invariant definition of the properties of the dividing shell. These require the vanishing of a linear combination of the expansion scalar and of the shear on the shell, as well as that of its flow derivative. In GPG coordinates, it summarizes as a vanishing of both first- and second-order flow derivatives of the areal radius.  % \\begin{comment} Such conditions being analog to a static shell, we established conditions for a vanishing expansion, in addition to the conditions already mentioned. As a consequence, the latter restriction would also require the vanishing of the radial gradient of the generalised velocity field. To maintain such expansion-free limit shell, so that the dividing shell would be static in the full sense, we obtain an additional equation of state that relates the gauge-invariant Hessian of the model with the quantity $\\rho+3p$ involved in the strong energy condition, on the dividing shell. Such condition being very restrictive, we settle the definition of the dividing shell as given by the first two conditions.  T \\end{comment} {}In order to illustrate our findings we have considered some simple examples of cosmological interest that provide realizations of our results. We have considered a $\\Lambda$-CDM model whereby we consider an LTB universe with dust and a cosmological constant. Notice that the simultaneous consideration of the latter two components yields a perfect-fluid model for the combined matter content. Moreover it can be seen as a simplified model of a dust universe within a cosmological setting coarsely provided by $\\Lambda$, which would then mimic the energy content of the background cosmological model with a rate of expansion much smaller than that of the pure dust collapse.  We have chosen initial conditions motivated by cosmological considerations and have discussed the existence of a dividing shell for those cases. We have also generalized a result of Ref.~\\citet{LaskyLun06b} for the case where a cosmological constant is present, which states that a perfect-fluid core embedded in a universe filled with a cosmological constant necessarily exhibits a dust transition between the perfect-fluid inner region and the outer vacuum region. This permits one to envisage this case as a generalization of the former $\\Lambda$-CDM examples.  Finally we should mention that a thorough discussion of global conditions represents a much harder problem, and remains an open problem since this involves the full characterization of a partial differential equations problem with boundary conditions in an open domain. % \\begin{comment} Even at the classical level of the Navier-Stokes equation this is still an open problem at present. Some progress has been made by us in that direction and will be addressed in future works, including the global behaviour of the inhomogeneous $\\Lambda$-CDM model with general initial conditions \\citet{LeDMM09a} and the proposal of a perfect-fluid ansatz \\citet{LeDMM09b}. \\end{comment} {}%%%%%%%%%%%%%%%%%%"
    },
    "0910/0910.0800_arXiv.txt": {
        "abstract": "{ The present constraints on fundamental constant variation ($\\alpha$ and $\\mu$) obtained in the radio range are reviewed, coming essentially from absorption lines in front of quasars of intermediate to high-z galaxies, through CO, HI, OH, HCO$^+$, HCN .. lines up to NH$_3$ and CII. With ALMA, the sensitivity to detect radio continuum sources in narrow bands will increase by an order of magnitude, and the expected progress is quantified. The relative advantage of the radio domain with respect to the optical one is emphasized. ",
        "introduction": "Molecular absorption lines in the radio domain are very narrow, and the high spectral resolution is well suited to the study of variation of constants, with space and time. Observations of absorption lines in the millimeter domain have been developped in the 1990's in both our own Milky Way (Lucas \\& Liszt 1994) and in very distant galaxies (Wiklind \\& Combes 1994). Up to now,  at intermediate and high redshift, between z$=$0.25 to z$=$0.89, only four absorption lines systems have been detected in the millimeter range and a fifth system at 0.765, at the OH-18cm lines (Kanekar et al 2005). Out of these 5 systems, 3 are intervening lensing galaxies (and the background quasar is multiply imaged), and 2 correspond to an absorption of the host (PKS1413+135, B3-1504+377). Provided that the background continuum source is strong enough, a large number of different molecules and rotational transitions can been observed (for an overview see Combes \\& Wiklind 1996; Wiklind \\& Combes 1994, 1995, 1996a,b, 1997). The detectability depends mainly on the background source intensity, and with the increased sensitivity of ALMA, we could expect to detect 30-100 times more sources. ",
        "conclusions": "Due to caveats (velocity offsets, variability, optical depths, hyperfine ratios, etc) a negative result is more easy to believe than a positive one!  There is many ways to improve the results: first obtaining more systems and more statistics, then improving the sensitivity and the signal-to-noise ratios, taking into account the simultaneity of observations for the lines to be compared. The number of continuum sources that will be use as targets for absorption studies will be one or two orders of magnitude larger than now with ALMA, although the millimeter domain is not favoured here by the K-correction.  It could then be interesting to search rest-frame 3mm systems at centimeter wavelengths, with EVLA, SKA precursors and in the future with SKA.   \\begin{figure*}[t!] \\centerline{ \\resizebox{11cm}{!}{\\includegraphics[angle=-90,clip=true]{sait_combes-f1.ps}} } \\centerline{ \\resizebox{11cm}{!}{\\includegraphics[angle=-90,clip=true]{sait_combes-f2.ps}} } \\caption{\\footnotesize Comparison between all observed constraints on variation for $\\alpha$ (top) and $\\mu$ (down). The data in the $\\alpha$ diagram are from OKLO (Damour \\& Dyson 1996), MMM (Murphy et al 2003, Srianand et al 2004), FeII (Quast et al 2004) and radio (Murphy et al 2001). The data in the $\\mu$ diagram are from 21-UV (Tzanavaris et al 2007), H$_2$ (Reinhold et al 2006, King et al 2008), and radio (Murphy et al 2008, Henkel et al 2009). } \\label{alpha} \\end{figure*}"
    },
    "0910/0910.2216_arXiv.txt": {
        "abstract": "High-resolution non-ideal magnetohydrodynamical simulations of the turbulent magnetized ISM, powered by supernovae types Ia and II at Galactic rate, including self-gravity and non-equilibriuim ionization (NEI), taking into account the time evolution of the ionization structure of H, He, C, N, O, Ne, Mg, Si, S and Fe, were carried out. These runs cover a wide range (from kpc to sub-parsec) of scales, providing resolution independent information on the injection scale, extended self-similarity and the fractal dmension of the most dissipative structures. ",
        "introduction": "In star forming disk galaxies, matter circulation between stars and the interstellar medium (ISM), in particular the energy input by supernovae, determines the dynamical and chemical evolution of the ISM, and hence of the galaxy as a whole. So far ISM models used radiative cooling assuming collisional ionization equilibrium (CIE), which is a good approximation providing the cooling time is much longer than the recombination times (e.g., Kafatos 1973). A condition verified for most of the ions for temperatures $> 10^{5.8}$ K. At lower temperatures departures from equilibrium are expected. While in CIE the ionization fractions depend only on the temperature and are sharply peaked, in NEI these same fractions depend on the dynamical and thermal history of the plasma. These departures affect the local cooling, which is a time-dependent process that controls the flow dynamics, feeding back to the thermal evolution by a change in the density and internal energy distribution, which in turn modifies the thermodynamic path of non-equilibrium cooling (Breitschwerdt \\& Schmutzler 1999). ",
        "conclusions": ""
    },
    "0910/0910.0025_arXiv.txt": {
        "abstract": "We report results from a $Chandra$ study of the central regions of the nearby, X-ray bright, Ophiuchus Cluster ($z=0.03$), the second-brightest cluster in the sky. Our study reveals a dramatic, close-up view of the stripping and potential destruction of a cool core within a rich cluster. The X-ray emission from the Ophiuchus Cluster core exhibits a comet-like morphology extending to the north, driven by merging activity, indicative of ram-pressure stripping caused by rapid motion through the ambient cluster gas.  A cold front at the southern edge implies a velocity of $1000\\pm200$ km s$^{-1}$ (M$\\sim$0.6). The X-ray emission from the cluster core is sharply peaked. As previously noted, the peak is offset by 4 arcsec ($\\sim2$ kpc) from the optical center of the associated cD galaxy, indicating that ram pressure has slowed the core, allowing the relatively collisionless stars and dark matter to carry on ahead. The cluster exhibits the strongest central temperature gradient of any massive cluster observed to date: the temperature rises from 0.7 keV within 1 kpc of the brightness peak, to 10 keV by 30 kpc. A strong metallicity gradient is also observed within the same region. This supports a picture in which the outer parts of the cool core have been stripped by ram-pressure due to its rapid motion. The cooling time of the innermost gas is very short, $\\sim5\\times10^7$ yrs. Within the central 10 kpc radius, multiple small-scale fronts and a complex thermodynamic structure are observed, indicating significant motions. These may be excited by the separation of the X-ray peak from the cD galaxy and its associated dark matter potential. Beyond the central 50 kpc, and out to a radius $\\sim150$ kpc, the cluster appears relatively isothermal and has near constant metallicity. The exception is a large, coherent ridge of enhanced metallicity observed to trail the cool core, and which is likely to have been stripped from it. ",
        "introduction": "\\begin{figure*} \\hspace{0.0cm} \\scalebox{0.43}{\\includegraphics{fig1a.ps}} \\hspace{0.9cm} \\scalebox{0.43}{\\includegraphics{fig1b.ps}} \\caption{Background subtracted, flat-fielded X-ray surface brightness maps in the $0.6-7.0$ keV band, smoothed with a 2 arcsec Gaussian filter. The three surface brightness discontinuities noted by Ascasibar \\& Markevitch (2006) are marked with arrows. (a) The central $8\\times8$ arcmin$^2$ region covered by ACIS chip 7. Note the comet-like extension to the north/northwest and the cold front observed $\\sim$40 kpc to the south of the X-ray peak (see also Ascasibar \\& Markevitch 2006). (b) The central $1.6\\times1.6$ arcmin$^2$ region of the cluster.  Two additional surface brightness discontinuities are observed.  One is 10 arcsec (5.5 kpc) to the northeast of the X-ray peak, the other is 3 arcsec (1.5 kpc) to the west (see also Ascasibar \\& Markevitch 2006). } \\label{fig:sb} \\end{figure*}  The Ophiuchus Cluster ($z=0.03$) is the second-brightest cluster of galaxies in the $2-10$ keV sky (Edge \\etal 1990; see also Ebeling \\etal 2002).  This extremely massive, nearby system lies behind the plane of our Galaxy and is highly obscured, with a line-of-sight equivalent hydrogen column density $N_H\\sim3\\times10^{21}$ atom cm$^{-2}$.  Using data from the Einstein Observatory and EXOSAT, Arnaud \\etal (1987) showed that the cluster has a high mean temperature, $kT=9.4^{+1.5}_{-1.2}$ keV (90 per cent confidence errors) and hosts a sharply-peaked cooling core. These findings were confirmed by later studies with ASCA (Matsuzawa \\etal 1996, Watanabe \\etal 2001) and $Chandra$ (Sun \\etal 2007).  Relatively cool, X-ray bright spectral components are associated with the cluster center, with a minimum measured temperature below 1 keV (P\\'erez-Torres \\etal 2009). The peak of the X-ray emission from the cluster is also offset from the X-ray centroid at larger scales (Arnaud \\etal 1987).  Ascasibar \\& Markevitch (2006) noted the presence of three surface brightness discontinuities in the central regions of the Ophiuchus Cluster (see also Fig. \\ref{fig:sb}), suggesting significant motion of the cool, dense core. Such fronts are common in clusters and are typically caused by merger events (\\eg Vikhlinin, Markevitch \\& Murray 2001; Reiprich \\etal 2001). However, many relatively relaxed clusters, such as Abell 1795 and 2029 (Fabian \\etal 2001; Markevitch, Vikhlinin \\& Mazzotta 2001; Clarke \\etal 2004), also exhibit surface brightness discontinuities near their X-ray surface brightness peaks. Simulations suggest that `sloshing' of the gas even in cool core clusters can lead to visible fronts up to several Gyrs after a relatively minor merger event (see Tittley \\& Henriksen 2005; Ascasibar \\& Markevitch 2006; see also Markevitch \\& Vikhlinin 2007 and references within). A cool core in motion within the ambient ICM will also experience ram pressure. The impact of this ram pressure is of particular interest for the survivability of cool cores following major mergers.  Due to its location and large angular size, optical studies of the Ophiuchus Cluster have proved challenging.  However, detailed spectroscopic studies suggest that the system is very massive, with a velocity dispersion $\\sigma_v=1050\\pm50$ km s$^{-1}$ (Wakamatsu et al. 2005), consistent with the measured X-ray temperature and luminosity (Edge \\etal 1990).  Clear substructure in the galaxy velocity histogram is observed, indicating that the cluster is merging (Wakamatsu \\etal 2005). The dominant cluster galaxy exhibits a substantial peculiar velocity, $v_p=$-650 km\\,s$^{-1}$, with respect to the cluster mean, $cz=9045\\pm30$ km\\,s$^{-1}$ (Wakamatsu \\etal 2005; Hasegawa \\etal 2000).  The presence of many nearby, smaller clusters and groups argues that Ophiuchus lies at the center of a supercluster (Wakamatsu \\etal 2005; see also Hasegawa \\etal 2000).  Govoni \\etal (2009; see also Murgia \\etal 2009) report the discovery of a low surface brightness, diffuse radio mini-halo surrounding the dominant cluster galaxy, 2MASX\\,J17122774-2322108. The central cD galaxy is also detected at 1.4 GHz with a flux density of 29.3 mJy in the NVSS survey (Condon \\etal 1998; see also Section \\ref{section:image}) with a $45\\times45$ arcsec$^2$ (FWHM) synthesized beam. There are a few other strong radio sources at the periphery of the cluster (see Johnston \\etal 1981). However, their association with the Ophiuchus Cluster is uncertain. A variety of authors have searched for non-thermal X-ray emission from the cluster in the hard X-ray band (see Nevalainen \\etal 2004, 2009; Eckert \\etal 2008; see also Dunn \\& Fabian 2006), although the results to date remain inconclusive (see Fujita \\etal 2008; Ajello \\etal 2009).  In this paper, we report detailed $Chandra$ observations of the core of Ophiuchus Cluster. Section 2 discusses the data reduction and analysis method. Section 3 describes the main results and Section 4 provides physical interpretation.  At a redshift of $z=0.03$, one arcsec corresponds to 0.60 kpc, assuming a $\\Lambda$CDM cosmology of $H_0=70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_M=0.3$, and $\\Omega_\\Lambda=0.7$. ",
        "conclusions": "The comet-like X-ray morphology, large-scale thermodynamic maps, and clear pressure jump at the location of the southern cold front (Figs. \\ref{fig:sb}a, \\ref{fig:lss}, \\ref{fig:zoom}), show that the cool, dense core of the cluster is moving rapidly through the large-scale hot, diffuse cluster medium. The velocity implied by the observed pressure jump at the southern cold front is $\\sim1000$ km s$^{-1}$.  The peculiar velocity of the cD galaxy of $\\sim650 $ km s$^{-1}$ with respect to the cluster mean (Wakamatsu \\etal 2005; Hasegawa \\etal 2000) also suggests significant, line-of-sight motion. Such speeds may indicate extreme `sloshing' of the cluster core (\\eg Markevitch \\etal 2001; Reiprich \\etal 2001; Tittley \\& Henricksen 2005; Ascasibar \\& Markevitch 2006; see also Markevitch \\& Vikhlinin 2007 and references within).  The X-ray emission associated with the cool core is very sharply peaked. The coolest ($kT \\sim 0.7$ keV), lowest-entropy gas lies within 1 kpc of the brightness peak. This material has a very short cooling time, with $t_{\\rm cool} \\sim4\\times10^7$yr, and plausibly represents the stripped interstellar medium of the dominant cluster galaxy.  The high relative velocity of the cool core argues that it was displaced from the base of the cluster potential well, presumably  by recent merger activity.  This is consistent with the slightly offset location of the X-ray peak with respect to the X-ray centroid at larger scales (Arnaud \\etal 1987). The case for a recent major merger event is supported by the significant substructure in the galaxy velocity histogram. Given the environment of the present-day cluster, which lies at the center of a supercluster and is surrounded by smaller galaxy groups and clusters, a recent merger event would not be surprising.  The degree to which cool cores are disrupted or destroyed during cluster mergers is an outstanding question in cluster physics, with important implications for cosmological studies (\\eg Burns \\etal 2008; Mantz \\etal 2009a,b; Ebeling \\etal 2009). The comet-like morphology of the cool core, and the remarkably steep central temperature gradient (rising from $\\sim 0.7$ keV to 10 keV within a 30 kpc radius) suggest that its outer regions have already been stripped by ram pressure, due to its rapid motion. Our data provide a dramatic, close-up view of the ongoing stripping and potential destruction of a cool core in a rich cluster and have clear implications for the survivability of such cool cores.  It will be important to examine, using constrained hydrodynamical simulations, the evolutionary path of such a cool core as it continues to be stripped and its motion slows.  At this point it is not clear how much of the original cool core will remain by the time the system returns to equilibrium.  The X-morphology of the inner 10 kpc region (Fig 1b), the separation between the X-ray and optical/near-IR brightness peaks (the X-ray emission trails the cD galaxy) and the offset between the X-ray peak and thermodynamic pressure maximum (the pressure maximum trails the X-ray brightness peak), all suggest additional, significant motions $within$ the cool core.  The sharp fronts noted in Fig. 1b suggest that significant rotational motion is present, and that the coolest, X-ray brightest gas is moving within the central potential. Such motion is not surprising, given the separation of X-ray peak from the stellar and, presumably, dark matter potentials of the cD galaxy (Fig. 2).\\footnote{More dramatic offsets between the X-ray emitting gas, and stellar and dark matter mass, have previously seen in massive, bimodal cluster mergers such as the Bullet Cluster (Clowe \\etal 2006; Bradac \\etal 2006) and MACS\\,J0025.4-1222 (Bradac \\etal 2008).} The sharp fronts seen in Fig 1b are likely to be sites of shearing instabilities, which could lead to rapid mixing of the coolest gas with its surrounding environment.  The ridges of enhanced metallicity in the cluster are particularly interesting. These large-scale, coherent structures are almost twice as metal rich as the surrounding gas. The prominent metallicity ridge 100 kpc to the north of the X-ray peak, and trailing in its wake, appears likely to have been stripped from the cool core during its motion. This ridge, and the partial ring-like ridge of high metallicity material at a radius of $\\sim$40 kpc, may initially have been driven outward by the central AGN, possibly even during the same outburst. The proximity of this ring to the southern cold front argues that this metal rich gas is currently being stripped from the cool core.  A qualitatively similar high-metallicity ridge is also seen in the central regions of the Perseus Cluster (Sanders \\etal 2005).  Govoni \\etal (2009) show that the cD galaxy is surrounded by a bright, diffuse mini-halo of synchrotron radio emission.  Despite the fact that a significant non-thermal particle population must be present to explain the radio emission, our study provides no evidence for non-thermal X-rays from the central regions of the cluster. High-resolution radio data for the cluster core are not currently available.  It will be important to obtain such data and examine the impact of the cluster dynamics on the properties of the central radio source, and the interaction between this source and its surrounding environment."
    },
    "0910/0910.5746_arXiv.txt": {
        "abstract": "In an effort to understand the origin of the Main Belt Comets (MBCs) 7968 Elst-Pizzaro, 118401, and  P/2005 U1, the dynamics of these three icy asteroids and a large number of hypothetical MBCs were studied. Results of extensive numerical integrations of these objects suggest that these MBCs were formed in-place through the collisional break up of a larger precursor body. Simulations point specifically to the Themis family of asteroids as the origin of these objects and rule out the possibility of a cometary origin (i.e. inward scattering of comets from outer solar system and their primordial capture in the asteroid belt). Results also indicate that while 7968 Elst-Pizzaro and 118401 maintain their orbits for 1 Gyr, P/2005 U1 diffuses chaotically in eccentricity and becomes unstable in $\\sim 20$ Myr. The latter suggest that this MBC used to be a member of the Themis family and is now escaping away. Numerical integrations of the orbits of hypothetical MBCs in the vicinity of the Themis family show a clustering of stable orbits (with eccentricities smaller than 0.2 and inclinations less than $25^\\circ$) suggesting that many more MBCs may exist in the vicinity of this family (although they might have not been activated yet). The details of the results of simulations and the constraints on the models of the formation and origins of MBCs are presented, and their implications for the detection of more of these objects are discussed. ",
        "introduction": "The detection of comet-like activities in three icy asteroids, 7968 Elst-Pizzaro, 118401, and P/2005 U1 (also known by their cometary indicators as 133P/Elst-Pizzaro, 1999 ${{\\rm RE}_{70}}$, and Read) by \\citet{Hsieh06} suggests the possibility of the existence of a new class of objects in the asteroid belt. Dubbed as Main Belt Comets, these objects have physical characteristics (e.g. comae and dusty tails) similar to those of comets, whereas dynamically, they resemble asteroids (their Tisserand numbers\\footnote{For a small object, such as an asteroid, that is subject to the gravitational attraction of a central star and the perturbation of a planetary body P, the quantity ${a_{\\rm P}}/a \\, + \\,2 {[(1-e^2)\\,a/{a_{\\rm P}}]^{1/2}}\\, \\cos i$ is defined as its Tisserand number. In this formula, $a$ is the semimajor axis of the object with respect to the star, $e$ is its orbital eccentricity, $i$ is its orbital inclination, and $a_{\\rm P}$ is the semimajor axis of the planet. In the solar system, the Tisserand number of a small body with respect to Jupiter can be used to determine the cometary or asteroidal nature of its orbit. In general, the Tisserand numbers of comets with respect to Jupiter are smaller than 3, whereas those of asteroids are mostly larger.} with respect to Jupiter are larger than 3.) As indicated by their orbital elements (Table 1), these objects are close to the outer region of the asteroid belt and within or in the proximity of the Themis family of asteroids (Fig.1). The dusty tails of these bodies persist for many weeks during their perihelion passages implying that tails of MBCs are not the ejections of impact-generated dust particles. An analysis of the dust trail of 7968 Elst-Pizzaro by \\citet{Hsieh04} suggests that small grains in the tail of this asteroid are dust particles that have been ejected from the surface of this body by the drag force of the gas produced by the sublimation of near-surface water ice.  The dual characteristic of 7968 Elst-Pizzaro, 118401, and P/2005 U1 (i.e. their apparent cometary activities and Tisserand numbers larger than 3) has raised many questions regarding the origin of these objects. On one hand, the physical appearance of these bodies may be taken as an evidence to argue that MBCs are comets that were scattered inward from the outer regions of the solar system and were captured in orbits in the asteroid belt during the early stages of the dynamical evolution of the solar system \\citep{Levison09}. On the other hand, the asteroidal-like Tisserand numbers of these bodies, combined with their proximity to the Themis family of asteroids can be used to argue the in-place formation of these objects in or around their current orbits. This paper examines these scenarios by studying the dynamics of these three MBCs and the constraints that the results may apply to the origin of these objects.  The currently known MBCs are a few kilometer in size. As indicated by Hsieh et al. (2004) and \\citet{Hsieh06}, the rate of the reduction of the sizes of these objects due to their cometary activities is approximately 1 meter per year. Similar to comets, the activities of MBCs are episodic and only intensify at their perihelion distances. As a result, such a rate of size-reduction implies that MBCs may not have long lifetimes and the current MBCs could not have started their activities too long ago. It also suggests that many icy asteroids might have had cometary activities in the past and are now inactive, and many more MBCs may exist in the outer region of the asteroid belt, which may be near their aphelia and have not started their activities yet. This paper discusses these issues by studying the dynamics of a large number of hypothetical MBCs for different values of their orbital parameters, and by identifying the regions of the parameter-space where these objects have higher probability of existence.  The outline of this paper is as follows. In section 2, the results of the numerical study of the dynamical properties of the three MBCs 7968 Elst-Pizzaro, 118401, and P/2005 U1 are presented. Section 3 has to do with the application of the results to the origin of these objects and their formation scenarios. Section 4 concludes this study by summarizing the results and discussing their implications for detections of more MBCs. ",
        "conclusions": ""
    },
    "0910/0910.3951_arXiv.txt": {
        "abstract": "The low mass X-ray binary 4U~2129+47 was discovered during a previous X-ray outburst phase and was classified as an accretion disk corona source. A $1\\%$ delay between two mid-eclipse epochs measured $\\sim 22$ days apart was reported from two {\\it XMM-Newton} observations taken in 2005, providing support to the previous suggestion that 4U~2129+47 might be in a hierarchical triple system. In this work we present timing and spectral analysis of three recent {\\it XMM-Newton} observations of 4U~2129+47, carried out between November 2007 and January 2008. We found that absent the two 2005 {\\it XMM-Newton} observations, all other observations are consistent with a linear ephemeris with a constant period of $18~857.63$\\,s; however, we confirm the time delay reported for the two 2005 {\\it XMM-Newton} observations. Compared to a {\\it Chandra} observation taken in 2000, these new observations also confirm the disappearance of the sinusoidal modulation of the lightcurve as reported from two 2005 {\\it XMM-Newton} observations. We further show that, compared to the {\\it Chandra} observation, all of the {\\it XMM-Newton} observations have 40\\% lower 0.5--2\\,keV absorbed fluxes, and the most recent {\\it XMM-Newton} observations have a combined 2--6\\,keV flux that is nearly 80\\% lower. Taken as a whole, the timing results support the hypothesis that the system is in a hierarchical triple system (with a third body period of at least 175 days). The spectral results raise the question of whether the drop in soft X-ray flux is solely attributable to the loss of the hard X-ray tail (which might be related to the loss of sinusoidal orbital modulation), or is indicative of further cooling of the quiescent neutron star after cessation of residual, low-level accretion. ",
        "introduction": "4U 2129+47 was discovered to be one of the Accretion Disk Corona (ADC) sources \\citep{for78}, which are believed to be near edge-on accreting systems since they have shown binary orbital modulation via broad, partial V-shape X-ray eclipses (\\citealt{tho79,mcc82}, hereafter MC82, \\citealt{wh82}); the eclipse width was $\\approx 0.2$ in phase and $\\approx 75\\%$ of the X-rays were occulted at the eclipse midpoint. The origins of ADCs are not fully understood, but they are typically associated with high accretion-rate systems, wherein we are only observing the small fraction of the system luminosity that is scattered into our line of sight. For the prototypical ADC X1822$-$371, models suggest that it is accreting at a near Eddington rate, with the radiation scattered into our line of sight having an equivalent isotropic luminosity on the order of 1\\% of the total luminosity (see \\citealt{par00,hei01,cot01}, and references therein).  In the early 1980s, observations of 4U 2129+47 showed that both its X-ray and optical lightcurves were modulated over a $5.24$~h period \\citep{tho79,ulm80,mcc82,wh82}. The discovery of a type-I X-ray burst led to the classification of 4U 2129+47 as a neutron star (NS) low mass X-ray binary (LMXB) system \\citep{gar87}, and the companion was suggested to be a late K or M spectral type star of $\\sim 0.6 M_{\\bigodot}$. The source distance was estimated to be $\\sim 1$--2\\,kpc \\citep{hor86}. Assuming isotropic emission, the X-ray luminosity would have then corresponded to $\\sim 5\\times10^{34}$~ergs$^{-1}$. Even if the luminosity were 100 times larger, the luminosity would have been somewhat smaller than expected for an ADC.  Since 1983, 4U 2129+47 has been in a quiescent state. Optical observations show a flat lightcurve without any evidence for orbital modulation between the years 1983 and 1987. Additionally, instead of an expected M- or K-type companion, the observed spectrum was compatible with a late type F8 IV star \\citep{kal88,che89}. The refined X-ray source position as determined by {\\it Chandra} turned out to be coincident with the F star to within $0.1^{''}$ (\\citealt{now02}, hereafter N02). The probability of a chance superposition is less than $10^{-3}$ (see also \\citealt{both08}), therefore, the hypothesis of a foreground star is unlikely. A $\\sim 40$\\,km~s$^{-1}$ shift in the mean radial velocity was derived from the F star spectrum, providing evidence for a dynamical interaction between the F star and 4U 2129+47 \\citep{cow90,both08}. The system was therefore suggested to be a hierarchical triple system, in which the F star is in a month-long orbit around the binary \\citep{gar89}. This hypothesis is tentatively confirmed with two {\\it XMM-Newton} observations separated by 22 days that showed deviations from a simple orbital ephemeris (\\citealt{boz07},hereafter B07). Assuming that the F star is part of the system, the source distance was revised to $\\sim6.3$~kpc \\citep{cow90}. This would imply that the peak luminosity observed prior to 1983 (assuming that the observed flux was $\\sim1$\\% of the average flux) was a substantial fraction of the Eddington luminosity. \\begin{figure*}[htpb] \\epsscale{1} \\plottwo{lc1.ps}{fit_fold.ps} \\caption {a) Left: Lightcurve from one of the 2007 {\\it XMM-Newton} observations, fit by a total eclipse with finite duration ingress and egress.  No sinusoidal orbital modulation was evident. The $90\\%$ confidence level upper limit of the modulation is $10\\%$. b) Right: Lightcurve of six eclipses folded as one. The mid-eclipse epochs are aligned and centered at 2250\\,s. \\label{fig:fig1}} \\end{figure*} The X-ray spectrum of 4U 2129+47, as determined with a {\\it Chandra} observation taken in 2000, was consistent with thermal emission plus a powerlaw hard tail (N02). The 2--8\\,keV flux was $\\approx 40\\%$ of the 0.5--2\\,keV absorbed flux. The contribution of the ``powerlaw'' hard tail to the 0.5--2\\,keV band was, of course, model dependent, but was consistent with being $\\approx 20\\%$ of the 0.5--2\\,keV flux, as we discuss below.  A sinusoidal orbital modulation (peak-to-peak amplitude of $\\pm30\\%$) was observed in the X-ray lightcurve in the same observation. The phase resolved spectra showed variation in the hydrogen column density.  However, in the {\\it XMM-Newton} observations (B07), the sinusoidal modulation was absent, while the flux of the powerlaw component was constrained to be less than $10\\%$ of the $0.2-10$ keV flux (B07).  Here we report on recent {\\it XMM-Newton} observations of 4U~2129+47. We combine analyses of these observations with reanalyses of the previous observations, and show that, absent the two 2005 {\\it XMM-Newton} observations (B07), the historical observations are consistent with a linear ephemeris with a constant period. We outline our data reduction procedure in \\S2 and present the timing and spectral analysis in \\S3 and \\S4 respectively. We summarize our conclusions in \\S5. ",
        "conclusions": "We reported on recent {\\it XMM-Newton} observations of 4U 2129+47 in a quiescent state. Given lower background flaring rates and relatively longer durations of these observations compared to the 2005 {\\it XMM-Newton} observations reported by \\citet{boz07}, we were able to make clear the following three points:  \\begin{enumerate} \\item The 0.5--2\\,keV absorbed X-ray flux of the source has been reduced by more than $\\approx 40\\%$ compared to the {\\it Chandra} observation. This was even true for the observations discussed by B07.  \\item The sinusoidal variation seen in the {\\it Chandra} observation ($\\pm30\\%$ peak-to-peak) is absent, or at least greatly diminished ($90\\%$ confidence level upper limits of $\\pm10\\%$ modulation). Additionally, the associated neutral column exhibits no orbital variability, and is consistent with the brightest/least absorbed orbital phases of the {\\it Chandra} observation.  \\item The powerlaw tail seen in the {\\it Chandra} observation is absent. A $\\Gamma=2$ powerlaw contributes less than $10\\%$ of the 0.5--2\\,keV flux, and is not evident in the $2-6$~keV band. The 2--6\\,keV flux is reduced by a factor of five compared to the {\\it Chandra} observation. \\end{enumerate}  With this further drop in the X-ray flux of the source, we may now truly be seeing 4U~2129+47 enter into a quiescent stage solely dominated by neutron star cooling, with little or no contribution from residual weak accretion. The spectrum of the source can be well explained by absorbed blackbody emission from the surface of a neutron star, with an emission radius of $\\sim1.8$~km. Likewise, more sophisticated neutron star atmosphere models also describe the spectra very well, with no need for an additional component.  As discussed by N02, the sinusoidal variability detected with the {\\it Chandra} observation could have been due to a neutral hydrogen column being raised by the interaction of the accretion stream with the outer edge of an accretion disk \\citep{par00,hei01}. The fact that the sinusoidal modulation was absent in the 2007, and likely also the 2005 {\\it XMM-Newton} observations could suggest that the system is in a lower accretion state, and might indicate that the geometry of the outer disk region has changed. It has been hypothesized that the hard tails of quiescent neutron stars originate from a pulsar wind, only able to turn on in quiescence, that is interacting with the accretion stream at large radius \\citep{cam98,cam00}. The vanishing of both the hard tail and the sinusoidal modulation in 4U~2129+47 could suggest that what we had observed in the {\\it Chandra} observation was indeed the interaction between the pulsar wind and the disk edge, perhaps as the last remnants of a disk accreted onto the neutron star. This is consistent with the fact that for these current {\\it XMM-Newton} observations the fitted neutral column is $\\approx 2$--$3\\times10^{21}\\,{\\rm cm}^{-2}$, which is comparable to the value measured for the least absorbed/peak of the {\\it Chandra} lightcurve. We note that a dominant nonthermal component has been reported in a ``recycled'' system: the binary radio millisecond pulsar PSR J0024$-$7204W. The X-ray variability observed in PSR J0024$-$7204W suggests that the hard X-ray emission is produced by interaction between the pulsar wind and matter from the secondary star, which occurs much closer to the companion star than the millisecond pulsar \\citep{bog05}.  The previously observed hard tail in 4U~2129+47 might have been a similar phenomenon.  What remains ambiguous is whether the associated drop in the 0.5--2\\,keV absorbed flux is solely related to the vanishing of the hard tail, or whether additional cooling of the neutron star has occurred, or continues to occur.  Although the soft X-ray flux has clearly declined from the {\\it Chandra} observations to the {\\it XMM-Newton} observations, there is no evidence for a decline between the 2005 and 2007/2008 observations.  The initial drop of the soft X-ray flux is consistent with solely the loss of the hard tail for some parameter regimes of this component; however, for the best-fit parameters, additional cooling would have had to occur.  We note that the expected luminosity of 4U~2129+47 due to neutron star cooling has recently been discussed by \\citet{heinke:08a}.  Even though 4U~2129+47 is among the brighter of the 23 quiescent sources presented by \\citet{heinke:08a}, the degree to which it is ``too faint'', and thus requires ``non-standard'' cooling (kaon cooling, pion cooling, etc.; see \\citealt{heinke:08a} and references therein) relies on the rather uncertain average heating history of 4U~2129+47, and the knowledge that the bulk of the observed soft X-ray emission was \\emph{not} due to residual accretion.  We note that any implied reduction of the soft X-ray emission of 4U~2129+47, or any continued cooling observed in the future, increases the need to invoke ``non-standard'' cooling mechanisms for this system, and further limits the parameter space for the long term average heating that would allow for ``standard'' cooling.  A question also arises as to whether or not the loss of the hard X-ray tail is in any way related to the unusual timing residuals associated with the 2005 {\\it XMM-Newton} observations (B07).  Strong deviations from a simple linear or quadratic ephemeris have also been noted for the bursting neutron star system EXO~0748$-$676 \\citep{wolff:09a}, with these changes possibly being correlated with the duration of the eclipse as viewed by the {\\it Rossi X-ray Timing Explorer} (RXTE). (Note that comparable eclipse duration changes, $<20$\\,s, observed in EXO~0748$-$676 are too small to have been seen in 4U~2129+47 with our observations.) \\citet{wolff:09a} hypothesize that these eclipse duration and ephemeris changes are related to magnetic activity within the secondary.  We note, however, that over a comparable five year span, the peak-to-peak residuals of EXO~0748$-$676 (70\\,s) are three to four times smaller than the peak-to-peak residuals seen in 4U~2129+47 (200--300\\,s).  Furthermore, the trend of ephemeris timing residuals of EXO~0748$-$676 are more persistently in one direction (although deviations do occur) than those that we show in Fig.~\\ref{fig:fig2} (compare to Fig.~4 of \\citealt{wolff:09a}).  In our work, we confirmed the strong deviations of two 2005 {XMM-Newton} eclipses from any linear or quadratic ephemeris, however, this deviation has only been observed once since 1982. Rather than presuming ephemeris residuals comparable to those of EXO~0748$-$676, given the optical observations \\citep[e.g.][]{gar87,both08}, we instead have adopted the hypothesis of a triple system.  Given the lack of constraints, however, we had to assume a sinusoidal function for simplicity. With this assumption, the chance of detecting a maximal time deviation and a minimal deviation should be equal. The rare presence of the strong deviation may be due to a large eccentricity of the third body orbit, so that the third star rarely interacts strongly with the inner binary system. The third body orbital period was estimated to be around 175 days, but can not be uniquely determined with the current observations. In order to confirm and accurately determine the orbital period of the third body, further observations are required. Such future observations could also address whether the soft X-ray flux of 4U 2129+47 has continued to drop, and whether the hard tail and sinusoidal modulation remain absent."
    },
    "0910/0910.3998_arXiv.txt": {
        "abstract": "We investigate the effect of two important, but oft neglected, factors which can affect the detectability of \\HI\\ 21-cm absorption in Mg{\\sc \\,ii} absorption systems: The effect of line-of-sight geometry on the coverage of the background radio flux and any possible correlation between the 21-cm line strength and the rest frame equivalent width of the Mg{\\sc \\,ii} 2796 \\AA\\ line, as is seen in the case of damped Lyman-$\\alpha$ absorption systems (DLAs). Regarding the former, while the observed detection rate at small angular diameter distance ratios ($DA_{\\rm abs}/DA_{\\rm QSO} > 0.8$) is a near certainty ($P>0.9$), for an unbiased sample, where either a detection or a non-detection are equally likely, at $DA_{\\rm abs}/DA_{\\rm QSO} \\geq 0.8$ the observed detection rate has only a probability of $P \\lapp10^{-15}$ of occuring by chance. This $\\gapp8\\sigma$ significance suggests that the mix of $DA_{\\rm abs}/DA_{\\rm QSO}$ values at $z_{\\rm abs}\\lapp1$ is correlated with the mix of detections and non-detections at low redshift, while the exclusively high values of the ratio ($DA_{\\rm abs}/DA_{\\rm QSO}\\sim1$) at $z_{\\rm abs}\\gapp1$ contribute to the low detection rates at high redshift.  In DLAs, the correlation between the 21-cm line strength ($\\int\\!\\tau\\,dv/N_{\\rm HI}$) and the Mg{\\sc \\,ii} equivalent width (${\\rm W}_{\\rm r}^{\\lambda2796}$) is dominated by the velocity spread of the 21-cm line. This has recently been shown not to hold for Mg{\\sc \\,ii} systems in general. However, we do find the significance of the correlation to increase when the Mg{\\sc \\,ii} absorbers with Mg{\\sc \\,i} 2852 \\AA\\ equivalent widths of ${\\rm W}_{\\rm r}^{\\lambda2852}>0.5$ \\AA\\ are added to the DLA sample. This turns out to be a sub-set of the parameter space where Mg{\\sc \\,ii} absorbers and DLAs overlap and the fraction of Mg{\\sc \\,ii} absorbers known to be DLAs rises to 50\\% \\citep{rtn05}. We therefore suggest that the width of the 21-cm line is correlated with W$_{\\rm r}^{\\lambda2796}$ for all systems likely to be DLAs and note a correlation between W$_{\\rm r}^{\\lambda2852}$ (Mg{\\sc \\,i}) and $N_{\\rm HI}$, which is not apparent for the singly ionised lines.  Furthermore, the 21-cm detection rate at $DA_{\\rm abs}/DA_{\\rm QSO} < 0.8$ rises to $\\gapp90$\\% for absorbers with ${\\rm W}_{\\rm r}^{\\lambda2852}>0.5$~\\AA\\ and large values of $DA_{\\rm abs}/DA_{\\rm QSO}$ may explain why the absorbers which have similar values of ${\\rm W}_{\\rm r}^{\\lambda2796}$ to the detections remain undetected. We do, however, also find the neutral hydrogen column densities of the non-detections to be significantly lower than those of the detections, which could also contribute to their weak absorption. Applying the $\\int\\!\\tau\\,dv/N_{\\rm HI}$--${\\rm W}_{\\rm r}^{\\lambda2796}$ correlation to yield column densities for the Mg{\\sc \\,ii} absorbers in which this is unmeasured, we find no evidence of a cosmological evolution in the neutral hydrogen column density in the absorbers searched for in 21-cm. ",
        "introduction": "\\label{intro}  Redshifted radio absorption lines can provide an excellent probe of the contents and nature of the early Universe, through surveys which are not subject to the same flux and magnitude limitations suffered by optical studies. In particular, with the \\HI\\ 21-cm line we can probe the evolution of large-scale structure, as well as measuring any putative variations in the values of the fundamental constants at large look-back times, to at least an order of magnitude the sensitivity provided by the best optical data. (\\citealt{cdk04} and references therein).  \\begin{table}% \\caption{Searches for intervening redshifted \\HI~21-cm absorption systems. The redshift range ($z_{\\rm abs}$) is given as well as the number of detections and non-detections ($n_{\\rm det}$ and $n_{\\rm non}$, respectively). \\label{comp}} \\begin{tabular}{@{}l c c c c  @{}} \\hline Reference      & Type    &  $z_{\\rm abs}$  & $n_{\\rm det}$ & $n_{\\rm non}$ \\\\ \\hline \\citet{br73} & DLA & 0.69 & 1 & 0\\\\ \\citet{rbb+76}    & DLA & 0.52 & 1 & 0 \\\\ \\citet{wd79} & DLA & 1.78 &  1 & 0 \\\\ \\citet{wbj81} & DLA & 1.94 & 1 & 0 \\\\ \\citet{bm83}$^{*}$ & sub-DLA & 0.44 & 1 & 0 \\\\ \\citet{bw83} & Mg{\\sc \\,ii} & 0.37--1.94 & 1 & 16 \\\\ \\citet{wbt+85}   & DLA  & 2.04 & 1 & 0 \\\\ \\citet{cry93}            & Lens  & 0.69 &  1 &  0\\\\ \\citet{cld+96} & DLA & 2.77--3.20 & 0 & 3\\\\ \\citet{dob96} & DLA  & 3.39 &  1 &  0\\\\ \\citet{lrj+96}           & Lens  & 0.19 &  1 & 0\\\\ \\citet{lsb+98}& DLA & 0.22--0.31 & 2 & 0\\\\ \\citet{cdn99}            & Lens & 0.89 & 1 &  0\\\\ \\citet{pvtc99} & DLA & 0.63 & 0 & 1\\\\ \\citet{ck00}$^{*}$ & DLA & 0.28--0.48 & 0 & 2\\\\ \\citet{lan00}$^{*}$ & Mg{\\sc \\,ii} &  0.21--0.96& 1  & 55\\\\ \\citet{lbs00}           & Ca{\\sc \\,ii} &0.09 & 1 & 0\\\\ \\citet{lb01}$^{*}$  & Mg{\\sc \\,ii} &0.44 & 1& 0\\\\ % \\citet{kc01a} & DLA & 0.25--0.56 & 2 & 1 \\\\ \\citet{kcsp01} & Mg{\\sc \\,ii} &0.10 & 0 & 1\\\\ \\citet{kc02} & DLA  & 0.42--3.18 & 1 & 9\\\\ \\citet{kb03}             & Lens & 0.76 & 1 &  0\\\\ \\citet{dgh+04}           & ---  & 0.78 &  1 &  --\\\\ \\citet{kse+06}           & DLA  & 2.35 &  1 &  0\\\\ \\citet{cdbw07}           & Lens  & 0.96 & 1 & -- \\\\ \\citet{ctp+07}           & DLA  & 0.66--2.71 & 1 &2  \\\\ \\citet{ykep07}           & DLA  & 2.29 &  1&  --\\\\ \\citet{gsp+09a}          & Mg{\\sc \\,ii} & 1.10--1.45 & 9 &  26 \\\\ \\citet{kpec09}$^{*}$& Mg{\\sc \\,ii} & 0.58--1.70 & 4 &  35\\\\ Zwaan et al. (in prep.)  & Mg{\\sc \\,ii} & $\\sim0.6$  & 2 & --\\\\ \\hline \\end{tabular} {$^{*}$Other detections reported, but which also appear in previous papers.} \\end{table} However redshifted \\HI\\ 21-cm absorbers are currently rare, with only 40 ``associated'' systems being detected in the hosts of radio galaxies and quasars (see table 1 of \\citealt{cww+08}) and another 40 ``intervening'' systems, which lie along the sight-lines to distant Quasi-Stellar Objects (QSOs), Table \\ref{comp}. In both cases, 21-cm absorption is predominantly detected at redshifts of $z\\lapp1$, which for the associated systems could be due to the excitation/ionisation caused by the proximity to the active nucleus, where optical surveys tend to select the most UV luminous sources at high redshift (\\citealt{cww+08}).%  For the intervening systems, many arise in known damped Lyman-$\\alpha$ absorption systems, which, at redshifts of $z_{\\rm abs}\\gapp1.7$, have the Lyman-$\\alpha$ line shifted into the optical band, allowing direct measurements of the neutral hydrogen column densities ($N_{\\rm HI}\\geq2\\times10^{20}$ \\scm, by definition). Non-detections can be thereby be attributed to high spin temperatures \\citep{kc02} and/or poor coverage of the background flux \\citep{cmp+03} in the high redshift systems (see Equ.~\\ref{enew}, Sect. \\ref{gen}).  Note, however, of the DLAs, the vast majority detected in 21-cm are also known Mg{\\sc \\,ii} absorbers, these traditionally being considered good candidates for the detection of 21-cm absorption at $z_{\\rm abs}\\lapp1.7$, where the Lyman-$\\alpha$ band is attenuated by the atmosphere. Two recent surveys of Mg{\\sc \\,ii} systems (at $1.10 < z_{\\rm abs} < 1.45$, \\citealt{gsp+09a} and $0.58 < z_{\\rm abs} < 1.70$, \\citealt{kpec09}), have found a total of 13 new 21-cm absorbers between them, significantly increasing the number known at $z_{\\rm abs}\\approx1$ (Table \\ref{comp}). In these works, various parameters (related to the equivalent widths of the singly ionised and neutral metal lines) are discussed, although the effects of geometry are generally ignored: \\citet{cw06} attribute the high 21-cm detection rate in DLAs identified through Mg{\\sc \\,ii}, cf. Lyman-$\\alpha$, absorption to the fact that the Mg{\\sc \\,ii} transition traces a lower redshift range ($0.2\\lapp z_{\\rm abs} \\lapp 2.2$, with ground-based telescopes, cf. $z_{\\rm abs}\\gapp1.7$). At redshifts of $z_{\\rm abs}\\gapp1.6$, the geometry of our flat expanding Universe, ensures that foreground absorbers are {\\em always} at larger angular diameter distances than the background QSOs, meaning that their effective coverage of the background radio continuum is generally reduced compared to the $z_{\\rm abs}\\lapp1.6$ absorbers (particularly those at $z_{\\rm abs}\\lapp1.0$). In this paper, we address this issue, investigating possible geometry effects on the Mg{\\sc \\,ii} sample as a whole, as well as discussing other possible effects on the detectability of 21-cm absorption. ",
        "conclusions": "Following the recent surveys of Mg{\\sc \\,ii} absorbers at $0.58 < z_{\\rm abs} < 1.70$ \\citep{gsp+09a,kpec09}, there has been a large increase in the number of intervening 21-cm absorption systems detected at these redshifts.  Both these works discuss various reasons regarding the detection of 21-cm absorption in Mg{\\sc \\,ii} systems, but these focus upon the various equivalent widths of the singly ionised and neutral metal lines, without considering the possible line-of-sight geometry effects, although these have been found to bias the detection rate for confirmed damped Lyman-$\\alpha$ absorption systems \\citep{cw06}. Considering this effect, we find for the  Mg{\\sc \\,ii} absorbers: \\begin{enumerate} \\item For the sample as a whole (or limiting this to strong Mg{\\sc \\,ii} absorbers, e.g. ${\\rm W}_{\\rm r}^{\\lambda2796} > 1$ \\AA), the mix of detections at low redshifts is what would be expected purely from chance, although the high non-detection rate at $z_{\\rm abs}\\gapp1$ is highly unlikely.  \\item When restrictions are applied to the sample, so that a large fraction of the Mg{\\sc \\,ii} absorbers are expected to consist of DLAs (according to \\citealt{rtn05}), detection rates of $\\approx100$\\% are reached at low redshift, while the low detection probabilities at  $z_{\\rm abs}\\gapp1$ remain significant. \\end{enumerate} Although other factors may contribute (such as lower column densities for the 21-cm non-detections, Sect. \\ref{rdng}), which may be unmeasurable (i.e. $T_{\\rm spin}/f$ in the absence of $N_{\\rm HI}$), it is apparent that the likelihood of detecting 21-cm absorption is correlated with the angular diameter distance ratio for all intervening absorption systems, where the absorbers at low redshift are subject to a range of $DA_{\\rm abs}/DA_{\\rm QSO}$ ratios and thus exhibit a mix of detections and non-detections, whereas those at $z_{\\rm abs}\\gapp1$ all have $DA_{\\rm abs}/DA_{\\rm QSO}\\sim1$, while tending to be non-detections.  Since a given absorber will occult the background flux much more effectively at a lower angular diameter distance ratio, this suggests a strong contribution by the covering factor, which is independent from the measured 21-cm line strength, from which the spin temperature/covering factor degeneracy cannot be broken (Sect.~\\ref{gen}): In some cases the covering factor has been estimated as the ratio of the compact unresolved component's flux to the total radio flux \\citep{bw83,klm+09}. However, even if high resolution radio images at the redshifted 21-cm frequencies are available, this method gives no information on the extent of the absorber (or how well it covers the emission). Furthermore, the high angular resolution images are of the continuum only and so do not give any information about the depth of the line when the extended continuum emission is resolved out. By using the angular diameter distance ratios, we completely circumvent these issues, although we can only discuss generalities which cannot determine values of  $T_{\\rm spin}$ or $f$ for individual systems. What we do find, however, from our statistically large sample (a total of 162 absorption systems), is that 21-cm detection rates are much lower when the absorber and background QSO are at similar angular diameter distances.  Saying this, the recent surveys \\citep{gsp+09a,kpec09} yield 13 new detections between them at $z_{\\rm abs}\\sim1$, where the angular diameter distance ratios are high ($DA_{\\rm abs}/DA_{\\rm QSO}\\gapp0.8$, Fig. \\ref{0.8_0.01}).  However, each of these surveys also yields a large number of non-detections, with both exhibiting a 25\\% detection rate\\footnote{The binomial probability of 26 or more non-detections out of a sample of 35 is $P(\\geq k/n)=0.0030$ (\\citealt{gsp+09a}, where at $1.10 < z_{\\rm abs} < 1.45$ all of the systems have $DA_{\\rm abs}/DA_{\\rm QSO}\\sim1$).}. This is about double the rate for the whole $DA_{\\rm DLA}/DA_{\\rm QSO}>0.8$ sample, but very close to that when restrictions are applied to include those Mg{\\sc \\,ii} absorbers which may also be DLAs (Table \\ref{stats}).  This is not surprising since the \\citet{kpec09} survey selects only those with ${\\rm W}_{\\rm r}^{\\lambda2796} > 0.5$ \\AA\\ (the minimum equivalent width of a DLA in the \\citealt{rtn05} sample -- see their figure 2) and \\citet{gsp+09a} select those with ${\\rm W}_{\\rm r}^{\\lambda2796} > 1$ \\AA, restricting this further.  Another possible factor affecting the detection of 21-cm absorption could be the correlation between the 21-cm line strength ($\\int\\!\\tau\\,dv/N_{\\rm HI}\\propto f/T_{\\rm spin}$) and the Mg{\\sc \\,ii} 2796 \\AA\\ equivalent width (${\\rm W}_{\\rm r}^{\\lambda2796}$, \\citealt{cw06}).  However, the same range of equivalent widths is also spanned by the non-detections, making this a questionable diagnostic with which to find 21-cm absorption systems. We do find, however, that eight of the ten non-detections have $DA_{\\rm abs}/DA_{\\rm QSO}>0.9$, whereas the 13 detections are much more spread out, spanning $0.4 < DA_{\\rm abs}/DA_{\\rm QSO} < 1$, which suggests that geometry effects may again be responsible for the detection rate.  If the $\\int\\!\\tau\\,dv/N_{\\rm HI}\\propto f/T_{\\rm spin}$--${\\rm W}_{\\rm r}^{\\lambda2796}$ correlation holds up, the fact that the non-detections increase the significance, suggests that these absorbers may be close to 21-cm detection. In any case, we suggest that ${\\rm W}_{\\rm r}^{\\lambda2796}$ in conjunction with the value of $DA_{\\rm abs}/DA_{\\rm QSO}$ may provide a diagnostic with which to find 21-cm absorption.  As discussed by \\citet{ctp+07}, the strongest contributer to this correlation appears to be the 21-cm profile width, which is not surprising as at W$_{\\rm r}^{\\lambda2796}\\gapp0.3$~\\AA, the range discussed here, the Mg{\\sc \\,ii} profile is dominated by velocity structure \\citep{ell06}. Including all of the Mg{\\sc \\,ii} absorbers degrades the correlation (as found by \\citealt{gsp+09a,kpec09}), although through the testing of many different equivalent widths and their ratios (e.g. ${\\rm W}_{\\rm r}^{\\lambda2796}$, ${\\rm W}_{\\rm r}^{\\lambda2852}$ \\& ${\\rm W}_{\\rm r}^{\\lambda2796}/{\\rm W}_{\\rm r}^{\\lambda2600}$), we find that the correlation increases for the Mg{\\sc \\,ii} absorbers for which the Mg{\\sc \\,i} 2852 \\AA\\ equivalent width is ${\\rm W}_{\\rm r}^{\\lambda2852}>0.5$ \\AA. At 13.60 eV, \\HI\\ has a similar ionisation potential to the singly ionised metal species (15.04 for Mg{\\sc \\,ii} \\& 16.18 eV for Fe{\\sc \\,ii}), although it is not surprising that the cooler component (i.e. the 21-cm) is better traced by the neutral metal species (i.e. Mg{\\sc \\,i}, which has a potential of 7.65 eV).  Furthermore, the fact that ${\\rm W}_{\\rm r}^{\\lambda2852}>0.5$ \\AA\\ is a sub-set of the ${\\rm W}_{\\rm r}^{\\lambda2796}/{\\rm W}_{\\rm r}^{\\lambda2600}$---$\\,{\\rm W}_{\\rm r}^{\\lambda2852}$ space which contains all of the DLAs of \\citet{rtn05}, suggests that the ${\\rm W}_{\\rm r}^{\\lambda2852}>0.5$ \\AA\\ selection is (mostly) adding further DLAs to the confirmed DLA sample \\citep{ctp+07}. That is, in general, the 21-cm profile width is correlated with W$_{\\rm r}^{\\lambda2796}$ when $N_{\\rm HI}\\geq2\\times10^{20}$ \\scm.  Investigating this further, in Fig. \\ref{pdlc} we show the total neutral hydrogen column density versus the Mg{\\sc \\,i} 2852 \\AA\\ equivalent width for the $z_{\\rm abs} < 1.7$ absorbers for which both measurements are available. \\begin{figure} \\centering \\includegraphics[angle=270,scale=0.75]{2852_N.eps} \\caption{The total neutral hydrogen column density versus the Mg{\\sc \\,i} 2852~\\AA\\ equivalent width for  $z_{\\rm abs} < 1.7$ absorbers \\citep{pdlc04}. The coloured symbols/arrows designate upper limits to ${\\rm W}_{\\rm r}^{\\lambda2852}$. The horizontal dashed line shows $N_{\\rm HI}=2\\times10^{20}$ \\scm, above which the absorbers are considered to be DLAs and the dotted line shows the least-squares fit to all of the points (taking account of the limits via {\\sc asurv}). } \\label{pdlc} \\end{figure} Although there is no clear partitioning, all (in an admittedly small sample) of those with ${\\rm W}_{\\rm r}^{\\lambda2852}>0.5$ \\AA\\ have $N_{\\rm HI}\\gapp10^{20}$ \\scm, and unlike \\citet{pdlc04}, who find no significant correlation between metal line equivalent widths (${\\rm W}_{\\rm r}^{\\lambda2796}$ \\& ${\\rm W}_{\\rm r}^{\\lambda2600}$) and $N_{\\rm HI}$, we find ${\\rm W}_{\\rm r}^{\\lambda2852}$ to correlate with $\\log_{10}N_{\\rm HI}$ at a $2.83\\sigma$ significance.  \\begin{figure} \\centering \\includegraphics[angle=270,scale=0.75]{tau-2852.eps} \\caption{As the left panel of Fig. \\ref{strength_W_ratio}, but with the velocity integrated optical depth of the 21-cm line on the ordinate and the equivalent width of the Mg{\\sc \\,i} 2852 \\AA\\ line on the abscissa. Note that for ${\\rm W}_{\\rm r}^{\\lambda2852}>0.5$ \\AA, {\\em all} of the 21-cm non-detections have $DA_{\\rm abs}/DA_{\\rm QSO}\\geq0.85$.} \\label{2852_ratio} \\end{figure} In Fig. \\ref{2852_ratio} we extend this to the sample searched for 21-cm absorption. Ideally, we would show the normalised line strength ($\\int\\!\\tau\\,dv/N_{\\rm HI}\\propto f/T_{\\rm spin}$, as in Fig. \\ref{strength_W_ratio}), however there are only 12 absorbers with measurements of both (comprising of seven 21-cm detections and five non-detections)\\footnote{Although these few points do give a $2.21\\sigma$ correlation between $\\int\\!\\tau\\,dv/N_{\\rm HI}$ and ${\\rm W}_{\\rm r}^{\\lambda2852}$.}. From this, we see a weak correlation between $\\int\\!\\tau\\,dv$ and ${\\rm W}_{\\rm r}^{\\lambda2852}$, although, as was noted in \\citet{ctp+07}, the neutral hydrogen column density is required in order to obtain the full picture (and $f/T_{\\rm spin}$). The absence of a measured $N_{\\rm HI}$ could be responsible for the fairly weak correlation, especially considering that the non-detections tend to have lower column densities (Sect. \\ref{rdng}).  Finally, for completeness, in Fig. \\ref{2796_ratio} we show $\\int\\!\\tau\\,dv$ versus the Mg{\\sc \\,ii} 2796 \\AA\\ equivalent width for the whole sample and see that, \\begin{figure} \\centering \\includegraphics[angle=270,scale=0.75]{tau-2796.eps} \\caption{As Fig. \\ref{2852_ratio}, but with the equivalent width of the Mg{\\sc \\,ii} 2796 \\AA\\ line on the abscissa.} \\label{2796_ratio} \\end{figure} although the correlation is negligible for the confirmed DLAs only \\citep{ctp+07}, it becomes quite significant over a sample which is an order of magnitude larger. Bear in mind, however, that the FWHMs for the 21-cm non-detections in Fig. \\ref{2852_ratio} and \\ref{2796_ratio} have been estimated from the FWHM--${\\rm W}_{\\rm r}^{\\lambda2796}-\\,$[M/H] correlations of the detections (Sect. \\ref{pw}), which will tend to increase any significance derived for these. This will be tempered somewhat by the optical depth limits for the non-detections, which are known. Therefore the significance of each of these correlations should be considered to lie somewhere between the two quoted values.  To summarise our findings regarding the detectability of 21-cm absorption in intervening Mg{\\sc \\,ii}  absorption systems: \\begin{itemize}  \\item On the basis of a 21-cm detection and non-detection being equally probable, at small angular diameter distance ratios ($DA_{\\rm abs}/DA_{\\rm QSO}\\lapp0.8$) the observed distribution is expected for the whole sample, whereas at high ratios the observed distribution is very improbable. This suggests that the angular diameter distance ($DA_{\\rm abs}$) plays a crucial r\\^{ole} in the detection of 21-cm absorption, which in turn suggests a strong covering factor dependence.  \\item Large angular diameter distance ratios may also explain why the non-detections span a similar range of Mg{\\sc \\,ii} 2796 \\AA\\ equivalent widths as the detections, which exhibit a 21-cm line strength--${\\rm W}_{\\rm r}^{\\lambda2796}$ correlation. Note also, that, in general, the non-detections have considerably lower neutral hydrogen column densities, which could contribute to making these harder to detect over the whole Mg{\\sc \\,ii} sample\\footnote{Although $N_{\\rm HI}$ is normalised out in our definition of line strength in the case of the $\\int\\!\\tau\\,dv/N_{\\rm HI}$---${\\rm W}_{\\rm r}^{\\lambda2796}$ correlation.}.  \\item Using this correlation to estimate the column densities for the absorbers in which these are unmeasured (particularly at $1.1\\lapp z_{\\rm abs}\\lapp1.6$), we find that the cosmological mass density of the neutral gas does not deviate significantly from $\\Omega_{\\rm neutral~gas}\\approx1\\times10^{-3}$ over the redshifts spanned by the optically selected absorbers ($0.1\\lapp z_{\\rm abs}\\lapp3.5$). This consistency with the general optically selected population supports the 21-cm line strength--${\\rm W}_{\\rm r}^{\\lambda2796}$ correlation. \\end{itemize} \\citet{ctp+07} suggested that the $\\int\\!\\tau\\,dv/N_{\\rm HI}$---${\\rm W}_{\\rm r}^{\\lambda2796}$ correlation was dominated by the 21-cm profile width in confirmed DLAs, although this does not appear to be the case for the Mg{\\sc \\,ii} sample in general \\citep{gsp+09a,kpec09}.  However, we find this significance to increase when only the absorbers with Mg{\\sc \\,i} 2852 \\AA\\ equivalent widths of ${\\rm W}_{\\rm r}^{\\lambda2852}>0.5$ \\AA\\ are added. This happens to be a sub-set of the Mg{\\sc \\,ii} absorbers with $1 < {\\rm W}_{\\rm r}^{\\lambda2796}/{\\rm W}_{\\rm r}^{\\lambda2600}<2$ {\\sc and} ${\\rm W}_{\\rm r}^{\\lambda2852}>0.1$ \\AA, $\\approx40\\%$ of which are known to be DLAs \\citep{rtn05}. The FWHM--${\\rm W}_{\\rm r}^{\\lambda2796}$ correlation does therefore appear to hold for absorbers in which the neutral hydrogen column density exceeds $N_{\\rm HI}\\sim10^{20}$ \\scm. Finally, we note a (somewhat scattered) correlation between the total neutral hydrogen column density and the equivalent width of the Mg{\\sc \\,i} 2852 \\AA\\ line, although there appears to be no such relation for the singly ionised metal species \\citep{pdlc04}."
    },
    "0910/0910.1353_arXiv.txt": {
        "abstract": "In scalar-tensor theories the scalar fields generically couple nontrivially to gravity. We study the observable properties of inflationary models with non-minimally coupled inflaton and Dirac-Born-Infeld (DBI) kinetic term. Within the assumptions of the priors of our Monte-Carlo simulations we find these models can generate new interesting observable signatures. Our discussion focuses on string theory inspired phenomenological models of relativistic D-brane inflation. While successful string theory constructions of ultra-violet DBI brane inflation remain elusive, we show that in suitable regions of the parameter space it is possible to use cosmological observables to probe the non-minimial coupling. Fortunately, the most observationally promising range of parameters include models yielding intermediate levels of non-gaussianity in the range consistent with WMAP 5-year data, and to be constrained further by the Planck satellite. ",
        "introduction": "There is substantial evidence indicating that the early Universe underwent a brief period of rapid, accelerated expansion, or inflation \\cite{Guth:1980zm,Linde:1981mu,Albrecht:1982wi}. In addition to solving conceptual and theoretical problems of the Standard Big Bang model (such as the monopole, flatness and horizon problems),  inflation  provides a quantum origin for the seeds of large scale structures in the Universe and produces theoretical predictions that are supported by detailed measurements of temperature anisotriopies in the cosmic microwave background radiation (CMB)~\\cite{Komatsu:2008hk}. But, despite its many phenomenological successes, the microphysical origin of inflation remains unclear; indeed, ever since its inception, it has been a challenge to derive successful inflationary models from fundamental particle physics~\\cite{Brandenberger:2007qi}. Fortunately, early work on D-brane inflation in string theory~\\cite{Dvali:1998pa,Alexander:2001ks,Dvali:2001fw,Burgess:2001fx,Brodie:2003qv,Kachru:2003sx}, has fruitfully led to promising realizations of the scenario~\\footnote{For recent reviews see, e.g. \\cite{McAllister:2007bg,Baumann:2009ni}.}.  Many inflationary models are driven, in part, by an inflaton field with non-canonical kinetic term, such as k-inflation \\cite{ArmendarizPicon:1999rj}, Ghost inflation \\cite{ArkaniHamed:2003uz}  and Dirac-Born-Infeld (DBI) inflation \\cite{Silverstein:2003hf}.  In these models the speed of sound can be significantly smaller than unity. Because cosmological perturbations travel at speeds less than that of light in these models, distinct observational signals are typically generated, including detectable levels of non-Gaussianity, e.g. \\cite{Garriga:1999vw,Silverstein:2003hf,Alishahiha:2004eh,Chen:2006nt}. As we demonstrate in this paper, non-Gaussianity is an additional observable providing new information that can be used to distinguish between certain inflationary models.  It is the nature of quantum field theory in curved spacetime that scalar fields generically couple non-minimally to gravity~\\cite{Birrell:1982ix}. In this paper we study effects of gravitational couplings in inflationary models with non-standard kinetic terms. In particular, we are interested in determining how such couplings alter cosmological observables. Our goal is to determine if it is possible to probe such couplings using measurements of non-Gaussianity together with measurements of other standard physical quantities. We ultimately focus on the theoretically well-motivated setting of a non-minimally coupled inflaton field in a DBI inflationary model; however, the methods we develop are applicable to more general settings as discussed below.  An outline of the paper is as follows: In \\S\\ref{motive}, we provide a brief introduction to gravitational couplings in theories with non-standard kinetic terms and DBI models of D-brane inflation. In \\S\\ref{analytic}, we derive analytic expectations for the observables of canonical and non-canonical models with conformal coupling. In \\S\\ref{discrim}, we discuss the prospects of observationally discriminating between pure DBI theories and DBI theory with a conformal coupling term. In \\S\\ref{monte}, we describe the priors and general methodology used in our Monte-Carlo simulations. In \\S\\ref{results}, we present our interpretation and an analytic understanding of the results of the analysis. Finally, we provide a summary of significant findings and further conclusions in \\S\\ref{conclude}. Technical calculations are relegated to a series of Appendices. ",
        "conclusions": "In this paper we studied the effects of a non-minimally coupled inflaton to the DBI action. We assumed an AdS form for the warp factor and studied a wide range of potentials and parameter values using a Monte-Carlo analysis. Our aim was to determine the prospects for observationally distinguishing non-minimal coupling from standard minimally coupled DBI by adopting a phenomenological approach rather than to conform with the latest model building technology in string theory. Our first finding was that: \\begin{itemize} \\item There is a strong degeneracy of observables in models with very large and models with undetectable levels of non-Gaussiantiy. \\end{itemize} This result is not particularly surprising given our analytic analysis of \\S \\ref{analytic}. However, for an observationally interesting intermediate level of non-Gaussianity with $f_{NL} \\in [20,40]$ we find that it is possible, in principle, to distinguish between minimally coupled and non-minimally coupled DBI models.~\\footnote{Although our priors are carefully chosen to minimize the differences between the models, our conclusions are dependent upon these priors and the assumptions discussed in detail elsewhere in this paper.} \\it Assuming a detection of intermediate levels of non-Gaussianity $f_{NL} \\in [20,40]$ and the value of the non-minimal coupling parameter $\\xi = 1/6$ \\rm we find:  \\begin{itemize} \\item For a given tensor-to-scalar ratio $r$ and $f_{NL}$, non-minimal DBI models tend to have a redder scalar spectral index $n_s$ than minimal DBI models. In other words, for a given scalar spectral index $n_s$ and tensor-to-scalar ratio $r$, $f_{NL}$ tends to be smaller for non-minimal DBI models than for minimal DBI models. \\end{itemize} Focusing our attention to the small tensor signal range we find: \\begin{itemize} \\item Any detection of a red spectral index $n_s$ below .97 combined with no detection of a tensor signal at Planck sensitivity effectively would rule out both non-minimal and minimal DBI inflation models. \\item Any detection of a red spectrum and no detection of a tensor signal at Planck sensitivity would rule out non-minimally coupled DBI models. \\end{itemize}  Finally, the goal of this paper was to initiate the study of observable effects of gravitational couplings in theories with non-standard kinetic terms. While we have focused on the particular example of conformal coupling and DBI kinetic term, our methods are applicable to more general systems represented by the action (\\ref{act}). We leave further study of such models for future research. \\newpage"
    },
    "0910/0910.2093_arXiv.txt": {
        "abstract": "We present results of the multiwavelength campaign on the TeV blazar Mkn~501 performed in 2006 July, including MAGIC for the very-high-energy (VHE) $\\gamma$-ray band and \\textit{Suzaku} for the X-ray band.  A VHE $\\gamma$-ray signal was clearly detected with an average flux above 200 GeV of $\\sim$20\\% of the Crab Nebula flux, which indicates a low state of source activity in this energy range. No significant variability has been found during the campaign. The VHE $\\gamma$-ray spectrum can be described by a simple power-law from 80 GeV to 2 TeV with a photon index of $2.8\\pm0.1$, which corresponds to one of the steepest photon indices observed in this energy range so far for this object. The X-ray spectrum covers a wide range from 0.6 to 40 keV, and is well described by a broken power law, with photon indices of $2.257\\pm0.004$ and $2.420\\pm0.012$ below and above the break energy of $3.24^{+0.13}_{-0.12}$ keV. No apparent high-energy cut off is seen above the break energy. Although an increase of the flux of about 50 \\% is observed in the X-ray band within the observation, the data indicate a consistently low state of activity for this source. Time-resolved spectra show an evidence for spectral hardening with a flux level. A homogeneous one-zone synchrotron self-Compton (SSC) model can adequately describe the spectral energy distribution (SED) from the X-ray to the VHE $\\gamma$-ray bands with a magnetic field intensity $B=0.313$ G and a Doppler beaming factor $\\delta = 20$, which are similar to the values in the past multiwavelength campaigns in high states. Based on our SSC parameters derived for the low state, we are able to reproduce the SED of the high state by just changing the Lorentz factor of the electrons corresponding to the break energy in the primary electron spectrum. This suggests that the variation of the injected electron population in the jet is responsible for the observed low-high state variation of the SED. ",
        "introduction": "Blazars are radio-loud active galactic nuclei (AGNs) viewed at small angles between the jet axis and our line of sight.  The bulk relativistic motion of the emitting plasma causes the radiation to be beamed in a forward direction, making the variability appear more rapid and the luminosity appear higher than in the rest frame due to the relativistic beaming effect~\\cite[e.g.,][]{Ree66, Ghi93}.  Blazars with only weak or entirely absent emission lines in the optical band are classified as BL Lacertae Objects (BL Lacs). Their spectral energy distributions (SEDs: in $\\nu F_{\\nu}$) are characterized by a two-bump structure~\\citep{Fos98}.  Since the discovery of the first extra-galactic TeV-photon emitter, Mkn421~\\citep{Pun92}, very-high-energy (VHE: $E>$ 80 GeV) emission has been confirmed in more than 20 BL Lac objects.  The SEDs of many of those BL Lacs show the peaks of the lower energy bump at UV to X-ray energies. These objects belong to the sub-class known as \"High-frequency peaked BL Lacs (HBLs)\"~\\citep{Pad95}.  Their non-thermal emission in this lower energy bump is commonly ascribed to synchrotron radiation from relativistic electrons, accelerated in the jet moving with relativistic bulk speed~\\citep[e.g.,][]{Ghi98}. The two-bump structure in SEDs of HBLs has been well explained by synchrotron self-Compton (SSC) models, where the target photons of inverse-Compton scattering for the higher energy bump are the synchrotron photons produced by the same electron population~\\citep[e.g.,][]{CandC02}.  In this model the high energy end of the electron spectrum is responsible for both X-ray and VHE $\\gamma$-ray emission.  The observed correlations of the X-ray and VHE $\\gamma$-ray fluxes during large flares of VHE $\\gamma$-ray emitting HBLs~\\cite[e.g.,][]{Tak96, Mar99, Kra01} provide strong experimental evidence for the SSC mechanism for HBLs. The target photons could also be produced in the accretion disk~\\citep{Der93} or in the broad-line region~\\citep[e.g.,][]{Sik94}. Alternatively, the high-energy emission can be also due to pions produced by accelerated protons and ions and subsequent pion decay~\\citep{Man93} or direct synchrotron emission from high-energy protons~\\citep{Aha00}.  The results of applying emission models to the data can provide information on physical parameters of the jet, such as the co-moving magnetic field, the population of the accelerated electrons, the Doppler boosting factor and the size of the emitting region. HBLs often show strong flux variability on time scales of less than 1 hr~\\citep{Gai96, Aha07, MAGIC501}. Hence, simultaneous multiwavelength observations over a wide energy range, covering in particular X-ray and VHE $\\gamma$-ray bands, are essential to study the physics of these high-energy radiation emitters.  Until a few years ago, simultaneous multiwavelength observations were only possible during flaring states due to the low sensitivity of the participating $\\gamma$-ray telescopes. In the VHE $\\gamma$-ray band, new generation of Imaging Atmospheric Cherenkov Telescopes (IACTs) such as MAGIC, H.E.S.S., and VERITAS, can access the energy range from below 100\\,GeV up to several TeV. These instruments also allow us to detect GeV-TeV $\\gamma$-ray signals within short observation times (several hours) even in quiescent source states. A comparison of emission model parameters for several different source states may allow us to reveal the origin of the jet activity.  The \\textit{Suzaku} X-ray satellite (see section~\\ref{sec:Suzaku}) features the most sensitive instruments among current X-ray detectors for time-resolved coverage of a wide X-ray energy band, from the soft to the hard X-ray energies, well beyond 10\\,keV. \\textit{Suzaku} already performed observations of several known TeV HBLs and successfully obtained time-resolved spectra up to hard X-ray energies~\\citep{Sato08, Tag08, Rei08}. The capability to perform successful multiwavelength observations for TeV-HBLs with MAGIC and \\textit{Suzaku} even in quiescent states has been shown in \\citet{Tag08} and \\citet{Rei08}.  Mkn~501 (redshift $z = 0.034$) is categorized as an HBL and was the second established TeV blazar~\\citep{Qui96}. In 1997 this source went into a state of surprisingly high activity. The detected flux was 10 times higher than that of the Crab Nebula in the VHE $\\gamma$-ray regime;  the high energy photons were observed up to $\\sim$20 TeV~\\citep{Aha01a}. At the same time, \\textit{Beppo}SAX also observed a flaring activity in the X-ray band~\\citep{Cat97, Pia98, Tav01}. The observed X-ray data showed an exceptionally hard spectrum with a synchrotron peak at (or above) $\\sim$ 100\\,keV. This represents a shift of at least two orders of magnitude with respect to previous observations. \\citet{Gli06} organized a long-term monitoring campaign in 2004, also covering the X-ray and TeV energy bands. They confirm the presence of a direct correlation between X-ray and VHE $\\gamma$-ray emission, which appears to be stronger when the source is brighter. In 2005, when MAGIC observed the object between May and July, the source flux varied by an order of magnitude. During the two most active nights, rapid VHE $\\gamma$-ray flux variability with a doubling time of a few minutes was observed~\\citep{MAGIC501}.  Several extensive SED studies based on multiwavelength observations of this object were reported~\\citep[e.g.,][]{Kat99, Kra00, Sam00, Tav01, Ghi02}. However, the multiwavelength observations that included VHE $\\gamma$-ray and X-ray instruments were only conducted during flaring states, but no simultaneous X-ray and VHE $\\gamma$-ray data were available for a low state of activity.  In 2006 July, a joint multiwavelength campaign between MAGIC, \\textit{Suzaku}, and the Kungliga Vetenskaplika Academy (KVA) optical telescope\\footnote{http://tur3.tur.iac.es/} was organized to observe Mkn~501. We succeeded in obtaining clear detections from the simultaneous observations both in the VHE $\\gamma$-ray band by MAGIC and in the X-ray band by \\textit{Suzaku}. In this paper, we report the observational results of this campaign.  The observations and data reduction for both instruments are briefly described in Section 2. In Section 3, we present the results of the measured light curves and spectra for each energy band.  In Section 4, we discuss the application of a simple one-zone SSC model to the SEDs obtained in this campaign and compare those to the historical data from flaring states. Finally, we summarize our results in Section 5. ",
        "conclusions": "Figure~\\ref{Mkn501SSC} shows the overall SEDs of Mkn~501 with data obtained during this multiwavelength campaign as well as some historical data.  The flux in the VHE $\\gamma$-ray band is corrected for absorption by the extra-galactic background light (EBL) using the \"low-IR\" model of \\citet{Kne04}. In our optical data, the host galaxy contribution $(12.0\\pm0.3)$ (mJy)~\\citep{Nil07} has already been subtracted.  Assuming a uniform injection of the electrons throughout a homogeneous emission region, we applied a one-zone SSC model, developed by \\cite{Tav98, Tav01}, to our campaign data for estimating physical parameters of the emitting region. A spherical shape (blob) of radius $R$ is adopted for the emission region, filled with a tangled magnetic field of intensity $B$.  The electron distribution is described by a smoothed broken power-law energy distribution with slopes $n_1$ from $\\gamma _{\\rm min}$ to the break energy $\\gamma _{\\rm b}$ and $n_2$ up to a limit of $\\gamma _{\\rm max}$ and with a normalization factor $K$. The relativistic effect is taken into account by the Doppler beaming factor $\\delta$.  The radius, $R$,  is selected to be $1.03\\times 10^{15}$ cm, corresponding to the value reported in \\cite{MAGIC501} for the SEDs observed in 2005. $\\gamma _{\\rm max}$ is set to be $10^7$ since no cut-off in the high energy ends of both the X-ray and the VHE $\\gamma$-ray spectra is detected. $\\gamma _{\\rm min}$ is fixed at 1 as a nominal value because this parameter does not affect the emission in the energy bands of our data. Since the data do not clearly indicate the positions of the synchrotron and inverse-Compton peaks, we cannot fully constrain the SSC parameters by precise fits to the SEDs. In addition, variability timescales cannot be determined for the data because of the quiescent state of the source.  Therefore, our aim is to reproduce the observed spectral behavior and correlations within a simple unifying picture rather than entering in details of the jet structure.  First, we applied the SSC model for the low state SED which is obtained during our multiwavelength campaign in 2006. The one-zone SSC model can reproduce the measured X-ray and VHE $\\gamma$-ray spectra in this low state of activity of the source as shown in Figure~\\ref{Mkn501SSC}. However, it is apparent that the model in Figure~\\ref{Mkn501SSC} underestimates the flux in a low energy range between radio and optical. Usually, homogeneous models cannot be used to explain the low frequency radio emission \\citep[see e.g.,][]{Pia98} due to efficient self absorption. In previous studies, \\citet{Kat01} used an inhomogeneous conical jet model proposed by~\\citet{Ghi85} to explain radiation from the low radio frequency up to the ultraviolet of Mkn~501. Here, along the same lines, we consider the energy range between X-ray and TeV for our one-zone SSC model.   On the basis of model parameters for this low state, we attempted to reproduce the SED obtained during the flare on 2005 June 30, using the same SSC model. There are no simultaneous X-ray data other than \\textit{RXTE}/ASM available at that time. The measured flux by the \\textit{RXTE}/ASM (pink triangle in Figure~\\ref{Mkn501SSC}) shows a compatible level in the X-ray spectrum to those taken by \\textit{BeppoSAX} (green dots in Figure~\\ref{Mkn501SSC}) on 1997 April 16. In addition, the VHE $\\gamma$-ray spectrum taken by MAGIC on 2005 June 30, was almost equivalent to the spectrum measured by CAT on 1997 April 16, as shown in Figure~\\ref{Mkn501VHEsp_all}. Therefore, we used this \\textit{BeppoSAX} spectrum as a guide for the X-ray spectrum during the VHE $\\gamma$-ray flare on 2005 June 30. With this, we can reproduce the SED in this high state just by changing $\\gamma _{\\rm b}$, the Lorentz factor of the electrons at the break energy in the electron spectrum. The SSC models for the low and the high states of Mkn~501 are represented by dashed lines in Figure~\\ref{Mkn501SSC}. The derived parameters for these SSC models are listed in Table~\\ref{Mkn501SSCp}.  In Table~\\ref{Mkn501SSC_comp} we compare our results to some of the previous SED studies based on SSC models for Mkn~501. All of those were derived from applying one-zone SSC models to actual observational data. Not all studies used simultaneous X-ray and VHE $\\gamma$-ray data. In fact, most simultaneous X-ray and VHE $\\gamma$-ray data were taken during the huge outbursts in 1997. Nevertheless, the SSC model parameters of $\\delta$ and $B$ from our multiwavelength campaign in 2006 indicate values similar to those of previous works for different flux states, apart from models of \\cite{Sam00} (lower $B$) for the medium flux state in 1998\\footnote{however, they did not take into account the absorption by EBL in the VHE $\\gamma$-ray data.}, and of \\cite{Kon03} (higher $\\delta$ and lower $B$) for the high flux state in 1997. Those parameters are also consistent with values in~\\citet{Bed99}, who constrain the parameters of the emission region based on the variability time scale during the 1997 April 15-16 flaring activity. Note again that only our work used simultaneous X-ray and VHE $\\gamma$-ray data taken in a low flux state.  \\cite{Tav01} attempted to model different emission states of Mkn~501 in 1997 and 1999 by mainly changing the break energy of electrons, slightly modifying its spectral slopes and number density, and by keeping other parameter unchanged. Also \\cite{Pia98} could reproduce different flux states by changing only the electron distribution, keeping the same values for other parameters. In these frameworks, the electron spectrum is the key component representing the different activity states of Mkn~501; especially, $\\gamma _{\\rm break}$ can play a major role there.  From the physical point of view, in the context of the widely discussed diffusive shock acceleration models~\\citep{Kir98, Hen99} the variations of $\\gamma _{\\rm break}$ (and the other parameters specifying the particle acceleration) could be explained by changes of the parameters determining the efficiency of the acceleration mechanisms (such as the parameters characterizing the turbulence). A deeper discussion of this point is clearly beyond the scope of this paper."
    },
    "0910/0910.3576_arXiv.txt": {
        "abstract": "Reliable atomic data have been computed for the spectral modeling of the nitrogen K lines, which may lead to useful astrophysical diagnostics. Data sets comprise valence and K-vacancy level energies, wavelengths, Einstein $A$-coefficients, radiative and Auger widths and K-edge photoionization cross sections. An important issue is the lack of measurements which are usually employed to fine-tune calculations so as to attain spectroscopic accuracy. In order to estimate data quality, several atomic structure codes are used and extensive comparisons with previous theoretical data have been carried out. In the calculation of K photoabsorption with the Breit--Pauli $R$-matrix method, both radiation and Auger damping, which cause the smearing of the K edge, are taken into account. This work is part of a wider project to compute atomic data in the X-ray regime to be included in the database of the popular {\\sc xstar} modeling code. ",
        "introduction": "The improved resolution and sensitivity of current satellite-borne X-ray observatories ({\\em Chandra} and {\\em XMM Newton}) are allowing the study of previously inaccessible weak spectral features of astrophysical interest. In the early stages of these missions, it was realized that absorption by near-neutral species was common, and the fact that all charge states (except the fully ionized) left identifiable imprints in the X-ray spectrum has proven to be a powerful diagnostic. Inner-shell absorption is important in the outflows of Seyfert galaxies in terms of both Fe L$\\alpha$ \\citep{sako01, behar01} and K$\\alpha$ lines of high-Z elements \\citep{behar02}, and also of elements in the first row of the periodic table such as oxygen \\citep{pradhan00, behar03, gar05}. Furthermore, inner-shell absorption of continuum X-rays from bright galactic sources is a useful diagnostic of the interstellar medium \\citep{yao09, kaastra09}.  Nitrogen K-shell absorption and emission are detected in X-ray spectra, mostly due to the H- and He-like charge states. For instance, observations of the emission lines of \\ion{N}{6} and \\ion{N}{7} in the ejecta of $\\eta$~Carinae by \\citet{leu03} have resulted in the lower bound ${\\rm N/O}>9$ for its nitrogen abundance. This result puts a constraint in the evolution of $\\eta$~Car and is a signature of CNO-cycle processing. The \\ion{N}{6} triplet is observed in the spectrum of the magnetic B star $\\beta$~Cep where it has been used to test magnetically confined wind shock models \\citep{fav09}. It is found that the plasma is not heated by magnetic reconnection and there is no evidence for an optically thick disk at the magnetic equator. Highly ionized emission lines of nitrogen have been observed by \\citet{miy08} on the north-eastern rim of the Cygnus Loop supernova remnant which can be used to determine nitrogen abundances, which turn out to be 23\\% solar. Seyfert galaxy outflows also can have super-solar N abundances \\citep{brinkman02, arav07}.  Narrow absorption K$\\alpha$ and K$\\beta$ lines of \\ion{N}{6} have been identified in the warm absorber of the MR 2251-178 quasar, which point to a complex velocity field with an outflow of ionized material \\citep{ram08}. K-shell absorption by Li-like \\ion{N}{5}, which has prominent UV lines, or by lower charge states at wavelengths $\\lambda > 29$~\\AA\\, is difficult to detect due to the reduced sensitivity of current X-ray instruments towards these longer wavelengths. Notable exceptions include : K$\\alpha$ absorption by \\ion{N}{5} at 29.42~\\AA\\ has been reported by \\citet{steenbrugge05} in the outflow of NGC~5548; and N absorption by a white dwarf outflow has been observed following the outburst of nova V4743 Sagittarii \\citep{ness03}. In the latter case, only the H- and He-lines are discussed, but lower charge states of N are clearly seen in the spectrum longward of 29~\\AA\\ \\citep[see Fig.~3b in][]{ness03}.  The current proliferation of X-ray spectra with high signal-to-noise ratio in astronomical archives makes the computation of nitrogen K-shell photoabsorption particularly timely. This is also an additional and important step in the ongoing effort to compute reliable atomic data by \\citet{gar05, pal08a, pal08b, wit09} for K-line analysis within the context of the {\\sc xstar} spectral modeling code \\citep{kal01}. Available spectroscopic data for K-vacancy levels of the N isonuclear sequence are notably scarce, mainly limited to \\ion{N}{5} and \\ion{N}{6} for which a few levels are listed in the NIST database \\citep{ral08} and four measured wavelengths have been reported by \\citet{bei99}. This shortage of reliable measurements precludes empirical corrections to calculated level energies. On the other hand, several relativistic methods have been previously used to generate atomic data for nitrogen K-vacancy states: the saddle-point complex-rotation method \\citep{dav89,chu90,shi01,lin01,lin02,zha05}; the Raleigh--Ritz variational method \\citep{hsu91, yan95,wan06}; and multiconfiguration Dirac--Fock (MCDF) approaches such as those by \\citet{hat83,har84,che86,che87,che88,che97}. Photoabsorption cross sections in the near K-edge region of N ions have been obtained by Hartree--Slater central-field computations \\citep{rei79}, where the resonance structure due to quasibound states as well as configuration correlations are neglected. The net effect of the resonance structure is to fill in the photoionization cross section below the inner-shell threshold altering the shape of the K edge.  In this paper we report on calculations of energy level structure and bound-bound and bound-free transition probabilities for the K-shell of nitrogen. The outline of the present report is as follows. The numerical methods are briefly described in \\S~2 while an analysis of the results based on comparisons with previous experimental and theoretical values is carried out in \\S~3. The two supplementary electronic tables are explained in \\S~4 while some conclusions are finally discussed in \\S~5. ",
        "conclusions": "Detailed calculations have been carried out on the atomic properties of K-vacancy states in ions of the nitrogen isonuclear sequence. Data sets containing energy levels, wavelengths, $A$-coefficients and radiative and Auger widths for K-vacancy levels have been computed with the atomic structure codes {\\sc hfr} and {\\sc autostructure}. High-energy photoionization and photoabsorption cross sections for members with electron occupancies $N\\geq 3$ have been calculated with the {\\sc bprm} and {\\sc hullac} codes.  Our best approximation (HF2) takes into account both core-relaxation effects and configuration interaction from the $n=3$ complex. By comparing results from different approximations with previous theoretical work and the few spectroscopic measurements available, we conclude that level energies and wavelengths for all the nitrogen ions considered in the present calculations can be quoted to be accurate to within 0.5~eV and 100 m\\AA, respectively. The accuracy of $A$-coefficients and radiative and Auger widths is estimated at approximately 20\\% for transitions neither affected by cancelation nor strong level admixture. Outstanding discrepancies are found with some MCDF data, in particular wavelengths and $A$-coefficients for the C-like ion which are believed to be due to numerical error by \\citet{che97}.  We have also presented detailed photoabsorption cross sections of nitrogen ions in the near K-threshold region. Due to the lack of previous experimental and theoretical results, we have also performed simpler calculations using the {\\sc hullac} code in order to check for consistency and to estimate accuracy. Comparison of {\\sc bprm} and {\\sc hullac} indicates that present K-threshold energies are accurate to within 1~eV. However, background cross sections are in better agreement with those computed by \\citet{rei79} in a central-field potential for all ions except the Li-like system. These photoabsorption cross sections and their structures are similar to those displayed by ions in the same isoelectronic sequence \\citep{wit09}, especially to the corresponding oxygen ions \\citep{gar05}.  The present atomic data sets are available on request and will be incorporated in the {\\sc xstar} modeling code in order to generate improved opacities in the nitrogen K-edge region, which will lead to useful astrophysical diagnostics such as those mentioned in \\S~1."
    },
    "0910/0910.4998_arXiv.txt": {
        "abstract": "We present a study of an extended \\lya\\ nebula located in a known overdensity at $z\\sim 2.38$.  The data include multiwavelength photometry covering the rest-frame spectral range from 0.1 to 250$\\mu$m, and deep optical spectra of the sources associated with the extended emission.  Two galaxies are associated with the \\lya\\ nebula. One of them is a dust enshrouded AGN, while the other is a powerful starburst, forming stars at $\\sim 600 M_{\\odot}$ yr$^{-1}$. We detect the \\heii\\ emission line at 1640~\\AA\\ in the spectrum of the obscured AGN, but detect no emission from other highly ionized metals (\\civ\\ or N~{\\small V}) as is expected from an AGN. One scenario that simultaneously reproduces the width of the detected emission lines, the lack of \\civ\\ emission, and the geometry of the emitting gas, is that the \\heii\\ and the \\lya\\ emission are the result of cooling gas that is being accreted on the dark matter halo of the two galaxies, Ly1 and Ly2. Given the complexity of the environment associated with our \\lya\\ nebula it is possible that various mechanisms of excitation are at work simultaneously. ",
        "introduction": "Searches for high redshift \\lya\\ emitters are now successfully finding large samples of normal star-forming galaxies at redshifts up to $z\\sim 6$.  These searches are also discovering rare, extremely powerful and spatially large \\lya\\ sources, often referred to as ``blobs'' in the literature \\citep{steidel2001}.  \\lya\\ nebulae are characterized by extended \\lya\\ emission on scales up to hundreds of kiloparsecs. The total integrated luminosity in nebulae can exceed $10^{43}$ \\es, that is, more than $10^{10}L_{\\odot}$ of energy is emitted in just the \\lya\\ line itself. These nebulae are similar both in energy output and morphology to the emission line nebulae associated with powerful high redshift radio galaxies, however no radio sources have been detected in their proximity. The number density of these sources is still poorly known, mainly because of the rarity of these sources; \\citet{yang2008} find only 4 \\lya\\ nebulae in a survey targeting $z=2.3$ over an area of $\\sim 5$ deg$^2$.  The \\lya\\ nebulae were first discovered by \\citet{steidel2001}, \\citep[see also ][]{francis1996,francis1997}, in a region characterized by an overdensity of galaxies by about a factor of 6 when compared to the field population.  Since then, most other nebulae have also been discovered in overdense regions of the Universe \\citep{prescott2008,matsuda2004}.  Apart from their extended size, what makes \\lya\\ nebulae peculiar compared to normal star forming \\lya\\ emitters is the lack of an obvious source of ionization for the gas.  Various mechanisms have been proposed to explain the observed emission, including photoionization by either powerful AGNs or starbursts \\citep[with the ionizing radiation escaping toward the nebula, rather than the observer;][]{steidel2001}, shocks produced by supernova driven winds \\citep{taniguchi2000,mori2004}, and cooling radiation from gas falling into a dark matter halo \\citep{fardal2001,yang2006}.  The latter model is of particular interest in light of recent theoretical studies showing that a substantial fraction of the in-falling gas does not shock heat to the virial temperature of the dark matter halo, but arrives at the central galaxy predominantly along filaments in the cosmic web \\citep{keres2005,dekel2006,dekel2008,brooks2008}. Such cold streams could be responsible for the observed \\lya\\ nebulae \\citep{yang2006}. This mode of gas accretion is predicted to be the way by which massive galaxies at $z\\sim 2$ can still acquire gas that can be turned into stars. However, an observational confirmation of this mechanism is still missing, although a few possible examples have been presented in the past few years \\citep{nilsson2006,smith2007,smith2008}.  Here we present an in--depth study of a complex \\lya\\ nebula at $z \\sim 2.38$ found to be associated with multiple 24$\\mu$m sources \\citep{palunas2004,colbert2006}. This \\lya\\ nebula (RA$=21:42:42.72$, DEC$=-44:30:00.0$) is located within a large structure of \\lya\\ emitting galaxies, in an $80\\times 80\\times 60$ comoving Mpc region surrounding the known $z=2.38$ group of galaxies J2143--4423 \\citep{francis1996}.  The structure is defined by $\\sim$ 30 spectroscopically confirmed \\lya\\ emitters \\citep{francis2004}, a space density a factor of $5.8\\pm 2.5$ greater than that found by field samples at similar redshifts.  The distribution of the \\lya\\ emitters is possibly filamentary, and contains several $5-10$ Mpc scale voids. This group of galaxies was first discovered as a cluster of Lyman limit absorption-line systems along the line of site of two $z>3$ QSOs \\citep[][]{francis2001,francis1993}.  The paper is organized as follows. Section~\\ref{sec:data} describes the imaging and spectroscopic data, Section~\\ref{sec:results} present the results. The origin of the galaxies associated with the nebula is studied in Section~\\ref{sec:sed}, while in Section~\\ref{sec:discussion} we discuss various plausible mechanisms for the gas ionization. Section~\\ref{sec:conclusions} presents our conclusions. Throughout the paper we assume $\\Omega_m = 0.3$, $\\Omega_m + \\Omega_{\\Lambda} = 1$, and $H_0 = 70$ km s$^{-1}$ Mpc$^{-1}$. All magnitudes are AB magnitudes \\citep{oke1974}, unless otherwise specified. ",
        "conclusions": "\\label{sec:conclusions} We have presented a multi-wavelength study of an extended \\lya\\ nebula identified in an over-dense filamentary structure at redshift $z=2.38$ \\citep{palunas2004}. Two sources (Ly1 and Ly2) of continuum are found to be physically associated with the extended \\lya\\ emission.  Both are powerful sources detected at 24$\\mu$m, but not detected at either $850\\mu$m or radio wavelengths, nor in the X-ray.  The power-law IR SED of Ly1 is consistent with that of a heavily reddened QSO, while Ly2 is a starburst galaxy, with properties similar to other powerful star-forming galaxies at $z\\sim 3$. The UV rest-frame spectra of these two objects confirm that the two galaxies belong to the $z=2.38$\\ structure.  They have only a small difference in velocity (a few hundred \\kms), so it is possible that we are witnessing an early stage of a merger between two massive galaxies.  The rest-frame UV spectrum observed at the position of Ly1 shows emission from \\heii\\ at 1640~\\AA, while we find no associated \\civ\\ or N~{\\small V} lines. Although it is possible that the observed emission line spectrum is generated in gas ionized by the continuum of the QSO, the the observed \\civ/\\heii\\ limit suggests that the \\lya\\ and the \\heii\\ radiation are the result of a different mechanism. The \\heii\\ line width argues against a WR stellar origin of the emission line. Extremely metal poor stars can be excluded as a source of ionization because of the presence of metal absorption lines in the UV spectrum of Ly2, and by the presence of dust inferred from the detection of both sources at 24$\\mu$m and the strong PAH lines in Ly2.  Another possibility that would be able to explain the observations is that the \\heii\\ and the \\lya\\ emission are produced by collisions in gas that is being accreted on the dark matter halo of the two galaxies, Ly1 and Ly2. Future rest-frame optical spectroscopy will help in the interpretation of the \\lya\\ cloud. If due to cold accretion, then the \\lya\\ is caused by collisional excitation rather than photoionization. In this case, the \\ha\\ line should be only about 2\\% of the \\lya\\ line \\citep{fardal2001}. Furthermore, the ratio between strong optical ([OIII], \\hb, \\ha, and [N~{\\small II}]) emission lines could be used to investigate the source of ionization in the nebula.  Given the complexity of the environment associated with our \\lya\\ nebula it is possible that the various mechanisms of excitation are at work simultaneously. This possibility has been discussed recently by \\citet{dijkstra2009}, who point out that the association of cold accretion with powerful starburst and/or AGN has to be expected, since ultimately these sources are triggered and fed by the gas infall.  Although theoretically the cold mode of accretion is expected to be the dominant mechanism for growth in massive galaxies at $z\\sim 2$ \\citep{dekel2008}, we still lack a direct observational confirmation of its existence. A common prediction among the different models of cold accretion is the unique filamentary morphology expected for the emitting gas. The typical filaments are predicted to have a surface brightness of about $10^{-19.5}$ erg s$^{-1}$ cm$^{-2}$ arcsec$^{-2}$ in both \\lya\\ and \\heii, and the cooling luminosity is expected to peak at redshifts between $\\sim 1$ and 2. At these redshifts, the \\lya\\ line is barely accessible from the ground and current space based facilities are not sensitive enough to reach the low surface brightness levels in reasonable exposure times. (Reaching a surface brightness of $10^{-19}$ erg s$^{-1}$ cm$^{-2}$ arcsec$^{-2}$ in the observed \\lya\\ line at $z=1.8$ with the Hubble Space Telescope Wide Field Camera 3 UVIS channel would require $\\sim 10^7$s). The proposed Advanced Technology Large-Aperture Space Telescope (ATLAS) could potentially detect cooling radiation in the redshift range (between $z=$1 and 2) where cold flows are expected to be the dominant mechanism for gas accretion in galaxies."
    },
    "0910/0910.1603_arXiv.txt": {
        "abstract": "We present \\xmmn\\ observations of four low-redshift Seyfert galaxies selected to have low host luminosities ($M_g > -20$ mag) and small stellar velocity dispersions (\\sigmastar\\ $< 45$ \\kms), which are among the smallest stellar velocity dispersions found in any active galaxies. These galaxies show weak or no broad optical emission lines and have likely black hole masses  $\\lesssim 10^{6}$~\\msun. Three out of four objects were detected with $> 3\\sigma$ significance in $\\sim 25$ ks exposures and two observations had high enough signal-to-noise ratios for rudimentary spectral analysis. We calculate hardness ratios ($-0.43$ to $0.01$) for the three detected objects and use them to estimate photon indices in the range of $\\Gamma = 1.1-1.8$. Relative to [\\ion{O}{3}], the type 2 objects are X-ray faint in comparison with Seyfert 1 galaxies, suggesting that the central engines are obscured. We estimate the intrinsic absorption of each object under the assumption that the [\\ion{O}{3}] emission line luminosities are correlated with the unabsorbed X-ray luminosity.  The results are consistent with moderate ($N_{\\rm H}~\\sim~10^{22}~{\\rm cm^{-2}}$) absorption over the Galactic values in three of the four objects, which might explain the non-detection of broad-line emission in optical spectra. One object in our sample, SDSS J110912.40+612346.7, is a near identical type 2 counterpart of the late-type Seyfert 1 galaxy NGC~4395. While the two objects have very similar [\\ion{O}{3}] luminosities, the type 2 object has an X-ray/[\\ion{O}{3}] flux ratio nearly an order of magnitude lower than NGC~4395. The most plausible explanation for this difference is absorption of the primary X-ray continuum of the type 2 object, providing an indication that obscuration-based unified models of active galaxies can apply even at the lowest luminosities seen among Seyfert nuclei, down to \\lbol\\ $\\sim 10^{40}-10^{41}$ \\ergs. ",
        "introduction": "Due to the advent of large area surveys in the past few years, extensive progress has been made in the search for low-mass active galactic nuclei (AGNs) with estimated black hole masses $\\mbh \\lesssim10^{6}~\\msun$. Through searches for galaxies with low stellar velocity dispersions or weak broad-line emission in the Sloan Digital Sky Survey (SDSS), the number of candidate broad-line type 1 \\citep{GH04, GH07c} and narrow-line type 2 \\citep{BGH08} low-mass AGNs has increased to number in the hundreds. Multi-wavelength studies (e.g. Gallo \\etal\\ 2008, Satyapal \\etal\\ 2008) have begun to provide new avenues for finding these low-mass and low-luminosity AGNs that would not typically be identified in the optical, allowing us to observe a larger portion of the total energy output. X-ray observations are of key importance for detecting the primary ionizing continuum of the AGN and determining the total luminosity as well as constraining the obscuration toward the central engine.  Current unification schemes explain the observational differences between type 1 and type 2 AGNs by the viewing angle from which we observe the central engine. For a type 2 object, an obscuring torus blocks the light coming from the innermost region, which contains the broad emission-line and X-ray-emitting regions. Therefore in the classical AGN unification picture \\citep{AM85}, a type 2 object will show the underlying properties of a type 1 object if the effects of the obscuration are removed. Hence, type 1 and 2 objects with comparable black hole masses and luminosities should also show similar optical narrow emission-line spectra, since the narrow-line emitting region remains largely unobscured. By systematically searching for galaxies of both types 1 and 2 with similar stellar velocity dispersions and narrow emission-line luminosities, one can develop comparable samples of type 1 and type 2 objects with similar black hole masses to test if the unified model is applicable across the full range of black hole masses found in AGNs, or if the lack of broad emission lines in low-mass type 2 objects results from changes in the structure of AGNs at low bolometric luminosities \\citep{Nicastro00, Laor, ES06} rather than due to absorption.  The two best studied AGNs with $\\mbh~<~10^{6}~\\msun$ are located in the late-type spiral NGC~4395 \\citep{FS89} and in the dwarf elliptical POX~52 \\citep{Kunth87, Barth04}, respectively. NGC~4395 has been shown to vary rapidly in the X-ray \\citep{Iwasawa00, Shih03, Moran05}, including dramatic changes in spectral slope ($\\Gamma~\\approx~0.6-1.7$, Moran \\etal\\ 2005) over a few years. POX~52 shows similar rapid variability, along with substantial changes in the absorbing column density in $< 1$~year \\citep{Thornton}. Additionally, \\xmmn\\ and \\chandra\\ have been used to investigate the X-ray properties of low-mass type 1 AGN samples \\citep{GH07a, Desroches, Mini} showing that they seem to be scaled down versions of their more massive counterparts with similar hardness ratios and photon indices. The X-ray properties of the type 2 counterparts of these AGNs have not previously been studied systematically.  We present \\xmmn\\ observations of four galaxies selected from the low-mass Seyfert 2 sample of \\citet{BGH08} to have the lowest stellar velocity dispersions in the sample in order to quantify the absorption and emission properties of this population. Our goal is to investigate whether obscuration can explain the lack of broad-line emission in low-mass Seyfert 2 galaxies. From measurements of X-ray luminosities and spectral fitting, we estimate intrinsic absorbing column densities, bolometric luminosities, and corresponding Eddington ratios in order to investigate the differences between type 1 and type 2 objects in this mass range. Throughout this paper we assume $H_{0}~=~70$~\\kms~Mpc\\per, $\\Omega_m$ = 0.3, and $\\Omega_\\Lambda = 0.7$. All estimates of Galactic foreground column densities are calculated based on Galactic \\ion{H}{1} maps \\citep{DL90, Kalberla05}, using the HEASARC online $N_H$ calculator\\footnotemark.  \\footnotetext{http://heasarc.gsfc.nasa.gov/cgi-bin/Tools/w3nh/w3nh.pl}  \\begin{deluxetable}{lllcc} \\tablecaption{Observation Details} \\tablehead{ \\colhead{Galaxy} & \\colhead{Obs. ID} & \\colhead{Obs. Date} & \\colhead{Exp. Time} & \\colhead{Cor. Exp.} \\\\ \\colhead{} & \\colhead{} & \\colhead{(UT)} & \\colhead{(s)} & \\colhead{Time (s)}} \\startdata SDSS J011905.14+003745.0  & 0400570301\t& 2006 Jul 26  & 26,739 & 18,040 \\\\ SDSS J103234.85+650227.9  & 0400570401\t& 2006 May 6   & 24,013 & 19,354 \\\\ SDSS J110912.40+612346.7  & 0400570201\t& 2006 Nov 25  & 23,613 & 23,613 \\\\ SDSS J144012.70+024743.5  & 0400570101\t& 2006 Aug 8   & 22,915 & 17,120 \\\\ \\enddata \\tablecomments{Cor. Exp. Time is the exposure time of each observation, corrected for the time lost due to background flares.} \\label{obsDate} \\end{deluxetable} ",
        "conclusions": "We find all four objects in the sample to be X-ray faint with X-ray luminosities approximately an order of magnitude lower than those seen in the \\citet{GH04} sample \\citep{GH07a}, but only two objects, \\one\\ and \\fourteen\\ show evidence of moderate to substantial absorption with estimated column densities of $N_{\\rm H} \\sim 10^{22}~\\mathrm{cm^{-2}}$. \\ten\\ shows little evidence of absorption which is consistent in the context of the unified model with the previous detection of broad emission lines from high-resolution optical spectra. It is unclear without further observations whether \\eleven\\ is truly lacking a broad-line region or if it is absorbed and the emission detected is due to a combination of X-ray binaries in the host galaxy and weak emission from the AGN, but given the low observed X-ray luminosity, we believe \\eleven\\ contains at least a moderate amount of absorption.  Two objects had high enough S/N ratios for spectral fitting. \\ten\\ has a spectrum similar to the partially absorbed spectrum of POX~52 discussed by \\citet{Thornton}, and is well fit with a high partial-covering fraction of 95\\%, in the $0.3~-~5.0$ keV range probed. The spectrum of \\fourteen\\ was well fit over an energy range of $0.3-1.0$ keV with a power law combined with a thermal plasma component and Galactic absorption and there was no clear evidence from the spectral fit for the high absorption predicted from the \\loiii-\\lX\\ correlation.  By comparing type 1 and 2 objects with similar masses, we have the opportunity to further investigate the absorbing properties of low-mass AGNs and whether or not line-of-sight absorption can explain the presence or lack of broad emission lines, or if other processes, such as accretion rate, play a role. If high S/N observations are obtained, X-ray spectra are an excellent tool that can be used to quantify the amount of absorption in an object. IR spectroscopy is another important tool used to investigate absorption and reprocessing of the AGN continuum. We plan to investigate both low-mass Seyfert 1 and 2 objects using {\\it Spitzer} spectra in order to better constrain the bolometric luminosities of these objects and further investigate the presence of an obscuring torus."
    },
    "0910/0910.3520_arXiv.txt": {
        "abstract": "We study the possibility to reconstruct primary mass composition with the use of combinations of basic shower characteristics, measured in hybrid experiments, such as depth of shower maximum from fluorescence side and signal in water Cherenkov tanks or in plastic scintillators from the ground side. To optimize discrimination performance of shower observables combinations we apply Fisher's discriminant analysis and give statistical estimates of separation of the obtained distributions on Fisher variables for proton and iron primaries. At the final stage we apply Multiparametric Topological Analysis to these distributions to extract composition from prepared mixtures with known fractions of showers from different primary particles. It is shown, that due to high sensitivity of water tanks to muons, combination of signal in them with $\\mathbf\\xmax$ looks especially promising for mass composition analysis, provided the energy is determined from longitudinal shower profile. ",
        "introduction": "\\label{sec:scan}  In the following we consider cases of Auger and TA to estimate the expected performance of the proposed mass reconstruction technique. We assume that primary energy can be estimated from the longitudinal shower profile and hence is practically primary mass independent. Briefly speaking, to enhance primary mass resolution of traditional \\xmax\\ parameter we suggest to use it in linear combinations with other basic shower parameters, such as signal in water tanks or particle density for Auger and TA correspondingly.  The data set used for the analysis was generated with COR\\-SI\\-KA~6.204~\\cite{corsika} (QGSJET~01~\\cite{qgsjet}/Gheisha~\\cite{gheisha}) and COR\\-SI\\-KA~6.735 (QGSJET~II~\\cite{qgsjetii}/Flu\\-ka2008.3~\\cite{fluka1}) packages and contains 1000 showers for every primary (p, O, Fe) and interaction model at 10~EeV and $37^\\circ$ zenith angle. All longitudinal showers characteristics and charged particles density were taken directly from CORSIKA output files. The calculations of the expected signal in Auger water tanks was performed according to the procedure described in~\\cite{billoir_sampl_2008,ave_munum_2007} with the use of the same GEANT~4 lookup tables as in~\\cite{ave_munum_2007}. In case of TA we use charged particles density at 1000~m from the axis ($D_{1000}$) as an example, for densities at another distances the consideration line would be the same. Finally let us note, that qualitatively results for both combinations of interaction models used in the study are very similar, so below we will mostly discuss only results for QGSJET~01/Gheisha, which provides some worse discrimination performance. To characterize the separation of distributions we will use the merit factor $ \\mathrm{MF}=|\\bar x_\\mathrm{Fe}- \\bar  x_\\mathrm{p}|/\\sqrt{\\sigma^2_\\mathrm{Fe}+\\sigma^2_\\mathrm{p}}, $ where $\\bar x_\\mathrm{p,\\,Fe}$ and $\\sigma_\\mathrm{p,\\,Fe}$ are distributions means and standard deviations correspondingly.  \\begin{figure} \\centering\\includegraphics[width=\\gsize\\textwidth]{icrc1511_fig1a.eps} \\centering\\includegraphics[width=\\gsize\\textwidth]{icrc1511_fig1b.eps} \\caption{Depth of shower maximum versus total ground plane signal and particle density for proton (red squares) and iron (blue crosses) showers at 1000 meters from the axis for QGSJET~01 model.} \\label{fig:sigxmax} \\end{figure}  In Fig.~\\ref{fig:sigxmax} we present scatter plots of total ground plane signal and charged particles density vs \\xmax\\ at 1000~m from the shower axis. Good separation of iron and proton showers in (\\smilla,\\,\\xmax) plot is both due to discrimination power of \\xmax\\ and to noticeable difference $\\sim13$\\% in the average total signals. In absolute units this difference is the same as the difference between muon signals, but the separation of the total signals (MF=1.4) is surely worse that of the muon ones (MF=2.5) due to the smearing effect of the electromagnetic component. Despite of change with the distance of the (electromagnetic/muon) signal ratio, a good separation of protons and iron nuclei is kept in a wide range from 700 to 1500 meters, since high sensitivity of Cherenkov tanks to the muon component results in different shifts of p and Fe populations along signal axis in \\xmax\\ vs total signal scatter plots. On the other hand, as expected, the separation between primaries in ($D_{1000}$,\\,\\xmax) plot is mostly due to discrimination power of \\xmax, since distributions on charged particles density for protons and iron nuclei largely overlap (MF=0.5). In this case one can think of searching for another discrimination parameters combinations, but as it will be shown below, ($D_{1000}$,\\,\\xmax) pair provides one of the best (among possible within TA conditions) discrimination resolutions. ",
        "conclusions": "The present study allows to develop a new approach to the mass composition analysis of hybrid data. We propose to use combinations of longitudinal and lateral parameters to achieve maximum primaries separation in multiparametric space. Further application of Fisher's method optimizes discrimination, reducing the problem to one-dimensional case and allowing for lower simulation statistics. At the last stage one can apply different algorithms to extract mass composition from distributions on Fisher variables, which are more mass sensitive in comparison with e.g. traditionally used \\xmax\\ alone. We applied for this purpose MTA technique to the samples with different primaries fractions and retrieved with very good accuracy nuclei abundances from [p,\\,Fe] and [p,\\,O] mixtures.  Regarding the choice of mass sensitive parameters, it was shown, that for Auger the best mass discrimination can be achieved if to use (\\xmax,\\,\\smilla) pair, provided the primary energy is estimated from the longitudinal shower profile. The charged particles density measured in TA in combination with depth of shower maximum also provides good discrimination of proton and iron showers, though in the real experimental conditions its sensitivity seems to be limited for proton-oxygen mixture case."
    },
    "0910/0910.4070_arXiv.txt": {
        "abstract": "We present a new analysis of the properties of star-forming cores in the Perseus molecular cloud, identified in SCUBA 850\\,\\micron\\ data originally presented by \\citet{hatchell05}. Our goal is to determine which core properties can be robustly identified and which depend on the extraction technique. Four regions in the cloud are examined: NGC~1333, IC348/HH211, L1448 and L1455. We identify clumps of dust emission using two popular automated algorithms, \\clfind\\ and \\gclumps, finding 85 and 122 clumps in total respectively. Using the catalogues of \\citet{hatchell07a}, we separate these clumps into starless, Class 0 and Class I cores. Some trends are true for both populations: clumps become increasingly elongated over time; clumps are consistent with constant surface brightness objects (i.e.\\ $M\\propto R^2$), with an average brightness $\\approx 4$--10 times larger than the surrounding molecular cloud; the clump mass distribution (CMD) resembles the stellar intial mass function, with a slope $\\alpha = -2.0 \\pm 0.1$ for \\clfind\\ and $\\alpha = -3.15 \\pm 0.08$ for \\gclumps, which straddle the Salpeter value ($\\alpha = -2.35$). The mass at which the slope shallows (similar for both algorithms at $M\\approx 6$\\,M$_\\odot$) implies a star-forming efficiency of between 10 and 20\\,per cent. Other trends reported elsewhere depend critically on the clump-finding technique: we find protostellar clumps are both smaller (for \\gclumps) and larger (for \\clfind) than their starless counterparts; the functional form, best-fitting to the CMD, is different for the two algorithms. The \\gclumps\\ CMD is best-fitted with a log-normal distribution, whereas a broken power law is best for \\clfind; the reported lack of massive starless cores in previous studies (e.g.\\ \\citealt{hatchell07a,hatchell08}) can be seen in the \\clfind\\ but not the \\gclumps\\ data. Our approach, exploiting two extraction techniques, highlights similarities and differences between the clump populations, illustrating the caution that must be exercised when comparing results from different studies and interpreting the properties of samples of continuum cores. ",
        "introduction": "Stars form out of gas in the densest areas of molecular clouds. Early studies of molecular cloud structure (e.g.\\ \\citealp{blitz80}) concluded that a division into discrete clumps of emission was the best description. Within such clumps, star formation occurs inside denser cores \\citep{myers83b}. A growing consensus now maintains that molecular clouds have a scale-free structure governed by turbulence (e.g.\\ \\citealp{elmegreen96,stutzki98,elmegreen04}), with clumps only an arbitrary categorization. We may reconcile these viewpoints to some extent; the clump population has a power-law spectrum of mass implying the overall cloud has a similar distribution (e.g.\\ \\citealp*{williams00}). Equally, the self-similarity of molecular clouds must break down where gravitational collapse becomes important.  This paper focusses on the decomposition of molecular clouds into \\emph{clumps} of emission. The utility of such a description depends on whether the located clumps accurately represent star-forming cores. Clumps are typically located in either dust continuum or spectral-line datasets, which both have their own advantages and drawbacks. Continuum observations select high-density cores, which tend to be self-gravitating, but associate more mass to clumps than exists in reality, since objects may be superposed along the line of sight (e.g.\\ \\citealp*{smith_clark08}). Clump-finding analyses on spectral-line cubes, with their added velocity dimension, should enable clumps to be more accurately assigned but are subject to larger uncertainties from the details of the radiative transfer than continuum emission. Ultimately, high signal-to-noise ratios are required in both cases for clumps to converge towards the underlying core population \\citep{ballesteros02}.  Nevertheless, many insights into star formation from clump populations remain compelling. \\citet*{motte98} first pointed out that the mass distribution of compact continuum clumps (the clump mass distribution; hereafter CMD) is remarkably similar to the initial mass function of stars (IMF). At high masses, the slope is consistent with a Salpeter power law ($\\mathrm{d}N/\\mathrm{d}M\\propto M^{-2.35}$, \\citealp{salpeter55}), considerably steeper than the power law for CO clumps  ($\\mathrm{d}N/\\mathrm{d}M\\propto M^{-1.6}$, e.g.\\ \\citealp{blitz93}). At lower masses, the CMD slope becomes shallower, around the point that samples start to suffer from incompleteness. This suggests that the mass of a star is established at its earliest phases (see also \\citealp*{alves07}) and would seem to rule out models where the shape of the IMF is formed later, through e.g.\\ dynamical interactions \\citep{bate05}. However, the mapping of the CMD on to the IMF is not straightforward and many evolutionary schemes would fit the present data (e.g.\\ \\citealp{swift08}).  Many wide-field surveys are about to or have just begun, which will locate and characterize thousands of star-forming cores from continuum data, e.g.\\ the James Clerk Maxwell Telescope (JCMT) and \\emph{Herschel} Gould Belt Surveys \\citep{gbs,andre05}. One of the key science drivers for these projects is the measurement of the CMD to high precision at low masses, into the brown-dwarf regime ($\\la 0.08$\\,\\msun). Thus, an examination of different source-finding techniques is particularly timely. Many different methods have been deployed in the literature to locate clumps in molecular-line and continuum data, from identifications by eye to automated algorithms. However, few studies \\emph{compare} the sources located with different techniques (a notable exception is \\citealt{schneider04}). In this paper, we closely compare the populations of continuum clumps found in the Perseus molecular cloud (hereafter simply Perseus) using the two most popular automated algorithms, \\clfind\\ \\citep*{williams94} and \\gclumps\\ \\citep{stutzki90}, to highlight their differences and biases. We re-examine SCUBA 850\\,\\micron\\ data, presented originally by \\citet{hatchell05}, in the four clusters of cores where we have complementary HARP spectral-line data \\citep*{paper1}, namely: NGC~1333, IC348/HH211 (IC348 for short), L1448 and L1455. Although only investigating a sub-set of the SCUBA data limits our sample size, these are the sources whose kinematics we will investigate subsequently and any CMDs from the spectral-line data will be directly comparable (Curtis \\& Richer, in prep.). The regions selected still encompass the majority of the SCUBA cores identified by \\citet{hatchell05,hatchell07a}, 58 out of 103; so we can extrapolate any conclusions to the entire cloud population. This work is divided into two principal parts. In Section \\ref{sec:clumppopulation}, we identify clumps using the two algorithms before matching them to a  catalogue of SCUBA cores, classified as protostellar or starless on the basis of their SEDs by \\citet{hatchell07a} (hereafter \\defcitealias{hatchell07a}{H07}\\citetalias{hatchell07a}). Second, Sections \\ref{sec:properties} and \\ref{sec:cmd} present an analysis of the physical properties of the cores before we summarize our conclusions in Section \\ref{sec:summary}.  \\subsection{Nomenclature}  Many different terminologies have been used to describe the hierarchical structure in molecular clouds; we follow \\citet{williams00}. Within molecular clouds, individual over-densities are termed \\emph{clumps}, these are the objects identified by automated algorithms and do not necessarily go on to form stars. Clumps may contain \\emph{cores}, which are the direct precursors of individual or multiple stars. Every clump that does not contain an embedded object is referred to as \\emph{starless}. Of these starless cores, only a subset, the \\emph{prestellar} cores (formerly pre-protostellar cores, \\citealt{wardthompson94})  will be gravitationally bound and thus go on to form stars. ",
        "conclusions": "\\label{sec:summary}  This paper detailed a new analysis of SCUBA 850\\,\\micron\\ data towards NGC~1333, IC348/HH211, L1448 and L1455 in the Perseus molecular cloud. We used two clump-finding algorithms, \\clfind\\ and \\gclumps, to decompose the dust column density into clumps of emission. \\clfind\\ located 85 clumps in total, whilst \\gclumps\\ found 122. Once identified we divided the clumps into protostellar (Class 0 and I) and starless sources by pairing them with classifications in the catalogues of \\citetalias{hatchell07a}.  This approach highlights trends in the clump properties that are highly or weakly dependent on the clump-finding algorithm. Conclusions which hold regardless of the technique include: \\begin{itemize} \\item \\textbf{Shapes}. Cores are slightly elongated with a tendency to become more elliptical over time as expected from collapse models. \\item \\textbf{Peak column densities}. Protostars and starless cores have similar average peak column densities. The upper half of the protostellar distribution arises from Class 0 protostars and the lower from Class Is. This is expected from the class definitions and supports the classifications of \\citetalias{hatchell07a}. Unsurprisingly, the masses of Class 0 clumps are also larger than Class Is. \\item \\textbf{Mass versus size}. Clumps are consistent with the $M \\propto R^2$ relation, having average surface brightnesses some 4-10 times larger than typical giant molecular clouds. \\item \\textbf{CMD}. The clump mass distribution resembles the IMF. The break mass is significantly higher than the average completeness mass and implies a star-forming efficiency of $\\sim 10$--20~per cent, probably reduced relative to embedded clusters due to multiplicity. \\item \\textbf{Regional variations}. The mass of clumps in each region follows: $M_\\mathrm{clump}(\\mathrm{L1448}) > M_\\mathrm{clump}(\\mathrm{NGC~1333}) > M_\\mathrm{clump}(\\mathrm{IC348})$. The high-mass slope of the {\\sc CMD} is steeper for IC348 than L1448 with NGC~1333 intermediate. \\item \\textbf{CMD temperature and lifetime dependence}. Using individual temperatures for each clump or weighting according to their free-fall timescale does not change the shape of the CMD. \\end{itemize}  Other inferences are dependent on the clump identification procedure: \\begin{itemize} \\item \\textbf{Sizes}. Protostars are either larger or smaller than prestellar cores with \\clfind\\ or \\gclumps\\ respectively. \\item \\textbf{Functional form of the CMD}. A log-normal CMD is best-fitting to the \\gclumps\\ data whereas a broken power law is best for \\clfind. In any case a broken-power-law fit has different slopes for each algorithm, with the \\gclumps\\ high-mass slope steeper at $\\alpha= -3.15\\pm 0.08$ than $\\alpha = -2.0\\pm 0.1$ for \\clfind. \\item \\textbf{A lack of massive starless cores?} A shortage of high-mass starless clumps relative to protostars was definitely present in the \\clfind\\ data (see also \\citetalias{hatchell07a}; \\citealp{hatchell08}). The same was not clear in the \\gclumps\\ data, where the overall spread in masses is smaller and the starless cores have larger sizes.  \\end{itemize}  \\gclumps\\ can only fit a strict shape so will not allocate flux to a clump at large distances from the peak unlike \\clfind. Therefore the corresponding \\gclumps\\ sizes are smaller. The clumps sizes in clustered regions such as NGC~1333 are smaller, as expected, for \\clfind\\ but not any different for \\gclumps, which is better at dealing with blended sources. The clump mass data are even more affected by selection effects. The slope of the starless and protostellar CMDs are very different for the two algorithms. Furthermore, the lack of high-mass starless cores, previously observed by \\citetalias{hatchell07a} and implied by the shallow prestellar mass function found by \\citet{enoch08} is not apparent in the \\gclumps\\ data. This would imply that either the shortage is not real or \\gclumps\\ does not allocate enough flux to high-mass clumps. Indeed, protostars which have more irregular, disrupted envelopes due to e.g.\\ outflows are perhaps worse affected by the regular profiles used by \\gclumps. This would cause a lack of high-mass protostellar clumps, which might explain the \\gclumps\\ results.  Fundamentally, we can never know the real population of clumps underlying any dataset. However, certain evolutionary trends \\emph{must} be present in any clump decomposition for it to be a believable representation of this true population. One clear example of this is the evolution of core size over time. Every simple model of collapse naturally has cores decreasing in size with time. This trend is \\emph{not} apparent for the \\clfind\\ clumps. In addition, particular \\clfind\\ clumps incorporate material that seems equally likely to go on to form part of their final mass or that of their neighbours. Therefore, treating the properties of a \\clfind\\ clump as representative of gravitationally-bound material in a single core is probably highly misleading. Of course, we may need to tweak the \\clfind\\ parameters, increasing the lowest contour level and decreasing the contour separation, but any method which necessarily prevents overlap in a highly-cluster region, such as Perseus, will struggle to find the true source population. Since one is normally interested in the material that will go on to form the final star, \\gclumps\\ would seem the more representative algorithm choice.  Analysing the clumps found by two of the most popular clump-finding routines side-by-side, underlines the caution that one must exercise when comparing results from distinct studies on similar datasets in the \\textit{same} region. Different methods can produce very different results from the same data, providing clump populations that are not directly comparable. Future continuum and molecular-line surveys will need to carefully design their source extraction procedures to ensure objects are found at the desired length-scale without simply using an automated algorithm blindly."
    },
    "0910/0910.0379_arXiv.txt": {
        "abstract": "{We present the results of new spectral diagnostic investigations applied to high-resolution long-slit spectra obtained with the Hubble Space Telescope Imaging Spectrograph (HST/STIS) of the jet from the T Tauri star RW~Aur.} {Our primary goal is to determine basic physical parameters (electron density $n_\\mathrm{e}$ and electron temperature $T_\\mathrm{e}$, hydrogen ionisation fraction $x_\\mathrm{e}$, total hydrogen density $n_\\mathrm{H}$, radial velocity $v_\\mathrm{r}$ and the mass outflow rate $\\dot M_\\mathrm{j}$) along both the red- and blueshifted lobes of the RW~Aur jet.} {The input dataset consists of seven long-slit spectra, of 0\\farcs1 spatial resolution, taken with the STIS slit parallel to the jet, and stepped across it. We use the Bacciotti \\& Eisl\\\"offel (1999) method to analyse the forbidden doublets [O\\,I]\\,$\\lambda\\lambda$6300,6363, [S\\,II]\\,$\\lambda\\lambda$\\,6716,6731, and [N\\,II]\\,$\\lambda\\lambda$\\,6548,6583\\,\\AA\\ to extract $n_\\mathrm{e}$, $T_\\mathrm{e}$, $x_\\mathrm{e}$, and $n_\\mathrm{H}$.} {We were able to extract the parameters as far as $3\\farcs9$ in the red- and $2\\farcs1$ in the blueshifted beam. The electron density at the base of both lobes is close to the critical density for [S\\,II] emission but then it decreases gradually with distance from the source. The range of electron temperatures derived for this  jet ($T_\\mathrm{e}$=$10^4$--$2\\times10^4$\\,K) is similar to those generally found in other outflows from young stars. The ionisation fraction $x_\\mathrm{e}$ varies between 0.04 and 0.4, increasing within the first few arcseconds and then decreasing in both lobes. The total hydrogen density, derived as  $n_\\mathrm{H} = n_\\mathrm{e} / x_\\mathrm{e}$, is on average $3.2\\times10^4\\mathrm{cm}^{-3}$ and shows a gradual decrease along the beam. Variations of the above quantities along the jet lobes appear to be correlated with the position of knots. Combining the derived parameters with $v_\\mathrm{r}$ measured from the HST spectra and other characteristics available for this jet, we estimate $\\dot M_\\mathrm{j}$ following two different procedures. The mass-outflow rate $\\dot M_\\mathrm{j}$ is moderate and \\textit{similar in the two lobes, despite the fact that the well-known asymmetry in the radial velocity persists close to the source.} Using the results of the BE diagnostics we find averages along the first 2\\farcs1 of both flows (a region presumably not yet affected by interaction with the jet environment) of $2.6\\times10^{-9}\\,M_\\odot$\\,yr$^{-1}$ for the red lobe and $2.0\\times10^{-9}\\,M_\\odot$\\,yr$^{-1}$ for the blueshifted flow, with an uncertainty of $\\pm \\log\\,M_\\odot$=1.6.} {{The fact that the derived mass outflow rate is similar in the two lobes appears to indicate that the central engine is constrained on the two sides of the system and that the observed asymmetries are due to environmental conditions. Possible suggestions for the origin of the differences are discussed. The RW~Aur jet appears to be the second densest outflow from a T Tauri star studied so far, but its other properties are quite similar to those found in other jets from young stars, suggesting that a common acceleration mechanism operates in these sources.}} ",
        "introduction": "\\label{intro}  Highly-collimated jets from young stellar objects (YSOs) have been thought for many years to be an essential ingredient of the star formation process \\citep{MF83}. Their origin, ejection, and collimation is, however, still a matter of debate \\citep[e.g.][]{Ray07}. Today, high-resolution spectroscopy from space and the ground allows us new insights into the physical conditions of these jets close to their source. This opens up new ways to compare observations with the different magneto-hydrodynamic (MHD) models proposed for the jet launch and acceleration \\citep[e.g.][]{Cam90,Koe00,Shu00,Matt03,Fer04}. During the past few years, high resolution studies of several jets from young stars on subarcsecond scales have become available -- e.g \\object{HH\\,30} \\citep{Bac99b,Har07}, \\object{DG~Tau} \\citep{Bac00,Lav00,Pyo03}, \\object{HL~Tau} \\citep{Pyo06}, \\object{RW~Aur} \\citep{Pyo06}, and \\object{LkH$\\alpha$~233} \\citep{Per07,Mel08}. The present study is a continuation of the spectroscopic analysis of RW~Aur performed with the Space Telescope Imaging Spectrograph (STIS) on-board the Hubble Space Telescope (HST), started by \\citet[][hereafter Paper~I]{Woi02}, and \\citet{Woi05}.  The classical T Tauri star RW~Aur is a double system in the Taurus-Auriga star forming region. The primary, RW~Aur~A (Sp. K1--K4), is one of the optically brightest T Tauri stars with $V=10\\fm1$. The secondary, RW~Aur~B ($V=12\\fm7$), is located at a projected separation of 1\\farcs2 and a position angle of $256^\\circ$ with respect to the primary. RW~Aur might be a multiple hierarchical system, although this is not clear. RW~Aur~A is a suspected spectroscopic binary \\citep{Pet01}. Another component, which was resolved very close to the secondary \\citep{Ghez93}, was probably a false detection \\citep{Whi01}.  A bipolar jet, with well-collimated blue- and redshifted lobes emanating from RW~Aur~A (\\object{HH\\,229}), was found by \\citet{Hir94}. Further work traced the emission out to 15\\arcsec\\ in the redshifted lobe \\citep{Dou00} and over 100\\arcsec\\ in the blueshifted beam \\citep{ME98}. Radial velocity measurements of several forbidden lines revealed an interesting  asymmetry: the radial velocity of the blueshifted lobe reaches about $-190$\\,km\\,s$^{-1}$, while a velocity of only $+100$\\,km\\,s$^{-1}$ is measured on the red side \\citep{Hir94}. Another unusual property of this jet is that for the first 10\\arcsec\\ the redshifted jet is brighter (in the red [S\\,II] doublet) than the blueshifted beam \\citep{ME98}, whereas in general blueshifted jets are brighter. Beyond this distance, the redshifted RW~Aur jet strongly drops in brightness, becoming much fainter than the blueshifted flow.  The physical parameters along the RW~Aur jet were first studied by \\citet{Bac96}, using spectra taken from the ground at moderate spatial and spectral resolution, and subsequently by \\citet{Dou00}, again from the ground but at higher spatial resolution using adaptive optics. These estimates are discussed in detail later in the present paper. A substantial advance in our understanding of this jet came, however, with the observations at high angular and spectral resolution performed with HST/STIS. In this case, seven long-slit spectra of the RW~Aur jet were taken at 0\\farcs1 spatial resolution, with the STIS slit parallel to the jet, and stepped every 0\\farcs7 across it. In \\citet{Woi02} (hereafter Paper I) we presented reconstructed velocity channel maps of the bipolar jet in the forbidden lines, giving  indications of the morphology and kinematics of the jet close to the star and at high resolution. In the second paper of the series, \\citet{Woi05}, we reported signatures of rotation in both jet lobes within the first 1\\farcs5 from the central source (about 200\\,AU at the distance of RW~Aur). Similar signatures of rotation have been discovered in other  jets from T Tauri stars \\citep{Bac02a, Cof04, Cof07}. Such observations support the magneto-centrifugal launching scenario for jets \\citep[cf][and references therein]{Fer06,Pud07}, which predicts a transport of angular momentum from the disk to the jet by the accretion/ejection `engine'. Recently, however, the disk surrounding RW~Aur has been found to rotate in the opposite sense of the bipolar jet \\citep{Cab06}. An explanation for this apparent discrepancy has not yet been found.  In this third paper of the series, we continue the analysis of our high-resolution HST/STIS spectra of the RW~Aur jet, to derive  the physical quantities in the emitting gas from the observed line ratios. Applying specialised diagnostic techniques we estimate the electron density and temperature, the ionisation fraction, the total hydrogen density and the mass outflow rate along the jet. Here we report on such results in detail, and we offer a new analysis of the spatio-kinematical structure of the bipolar jet extracted from the high-resolution position-velocity maps.  The paper is structured as follows: the observations and details of data reduction and analysis are described in Sect.~\\ref{obs}. In Section~\\ref{results} we illustrate the estimates of the physical quantities along both lobes of the jet. Then in Section~\\ref{disc} we contrast the derived jet properties of the opposing lobes, and we compare the estimates with values derived previously and in other T Tauri jets. Finally, we summarize our conclusions in Sect.~\\ref{concl}. ",
        "conclusions": "\\label{concl}  In this paper we analyse in detail the first few arcseconds of both lobes of the RW~Aur jet, using long-slit high-resolution HST/STIS spectra. The physical properties of the outflow are derived applying the so-called BE technique to the spectra, described, e.g., in \\citetalias{Bac99a} and in \\citet{Pod06}. Our results  can be summarised as follows:  \\begin{enumerate}  \\item The distribution of forbidden emission along  both outflows reveals several local maxima identified as jet knots. We see seven emission knots in the redshifted lobe within 3\\farcs8 from the source, and at least two knots in the blueshifted lobe within 2\\farcs1. The distribution of emission within the knots in the blue lobe appear to be flatter than in the redshifted section.  \\item The electron density $n_\\mathrm{e}$ in both outflows is close to the high density limit ($\\log(n_\\mathrm{e})_\\mathrm{crit}=4.39$) in the region near the driving source and decreases gradually with distance from the star. The electron temperature $T_\\mathrm{e}$ in the redshifted lobe decreases steeply within the first 0\\farcs3 from $1.5\\times10^4$\\,K to $10^4$\\,K and then varies around this value. The same trend is observed in the first few positions in the blueshifted lobe, but then $T_\\mathrm{e}$ reaches about $10^4$\\,K and subsequently shows higher values with a maximum of about $2\\times10^4$\\,K.  \\item We found that the gas in the red lobe has an ionisation fraction $x_\\mathrm{e}$ ranging between 0.02 and 0.2, with an average of 0.08; in the blue lobe this parameter is higher and varies between 0.1 and 0.4 (average 0.23). $x_\\mathrm{e}$ increases in the first section of both lobes, reaches a maximum at 1--2 arcseconds from the source and then starts to decrease. This behaviour appears to be a common feature also for other jets from young stars. Our results for $x_\\mathrm{e}$ for the red lobe agree within a factor of two with those of previous studies \\citep{Dou02,Bac96}.  \\item The hydrogen density $n_\\mathrm{H}$ in the brighter redshifted lobe appears to be twice as high (mean value $n_\\mathrm{H}\\approx 6.5\\times10^{4}$cm$^{-3}$) than in the fainter blueshifted section, where an average value of $n_\\mathrm{H}\\approx 2.3\\times10^{4}$ cm$^{-3}$ is found. RW~Aur jet is found to be the second densest \\citep[after DG~Tau, ][]{Lav00} of the T Tauri jets investigated so far.  \\item The kinematical properties of the flow have been fully re-analysed at high resolution and illustrated in accurate graphic detail. We confirm that the asymmetry in radial velocity ($-190$\\,km\\,s$^{-1}$ and  $+100$\\,km\\,s$^{-1}$ in the blue and red lobes, respectively) persists in all emission lines  and down to the source, and that the flow presents indications of rotation around its axis, in the same sense as determined in \\cite{Woi05}. The pattern of proper motions of the knots has been updated with the data presented here.  \\item Strong variations of $n_\\mathrm{e}, x_\\mathrm{e}$ and $T_\\mathrm{e}$ are observed close to knots J3--J4 in the red lobe and A12 in the blue lobe. The radial velocity variations are very limited and reach a maximum of 20\\% of the average at knots J3--J4 in the red lobe. Variations of the physical parameters at the other knot positions are suggested but are much less prominent.  \\item The average  mass flux $\\dot M_\\mathrm{j}$ calculated from the forbidden line ratios is similar in the first 2\\farcs1 of both lobes: $2.6\\times10^{-9}\\,M_\\odot$\\,yr$^{-1}$ for the red- and $2.0\\times10^{-9}\\,M_\\odot$\\,yr$^{-1}$ for the blueshifted jet. The $\\dot M_\\mathrm{j}/\\dot M_\\mathrm{acc}$ ratio is rather low and we estimate it to be between  0.003 and 0.012, depending on the adopted values for the mass accretion rate. This ratio, however, can still be accommodated by magneto-hydrodynamic models having a large value of the lever arm. We note however that the ratio may also be influenced by non-stationary accretion.  \\item The overall trend of the physical parameters along the first few arcseconds of the RW~Aur jet is similar to that of HH\\,30 and DG~Tau \\citep{Bac99b,Bac00, Bac02a, Har07}, which have been investigated with similar resolution. Analogies in the  properties of the RW~Aur, HH\\,30, DG~Tau jets in regions close to central star can reflect analogies in the mechanisms operating in that region, suggesting the same engine is  accelerating the jets in the T Tauri stars with outflows.  \\item Our study of the RW~Aur jet indicates for the first time that, despite the detected marked asymmetries in physical and kinematic properties between the two lobes, {\\em the  mass outflow rates in the two lobes are similar. This appears to indicate that the central engine has constraining symmetries on both sides of the system, and that the observed asymmetries are probably due to different environmental conditions.} We discuss the likelihood that these are  related to a large-scale inhomogeneity in the distribution of the local magnetic fields.  \\end{enumerate}  In conclusion, our HST/STIS study of the physical parameters of the bipolar jet from RW~Aur has highlighted a number of properties that appear to be common to many objects of the same class when observed close to their source. That said, it reveals that peculiar elements of this and other jets, like the observed asymmetries in the two opposite lobes, remain unexplained even after analysis conducted at sub-arcsecond resolution. Nevertheless, this work shows that an accurate investigation close to the source of dynamically relevant properties, like the mass outflow rate, that present no asymmetry, can constrain models. Obviously whether mass outflows rate is balanced in other asymmetrical bipolar jets should be investigated. The combined efforts of future theoretical studies, and soon-to-become available interferometric observations at milliarcsecond resolution are expected to give useful signposts to an understanding of these phenomena."
    },
    "0910/0910.0465_arXiv.txt": {
        "abstract": " ",
        "introduction": "Counting the number of planetary nebulae (PNe) in the Galaxy is not just a bookkeeping exercise. The PN phase enriches the ISM with nitrogen, carbon, helium, and dust, important components in the formation of future generations of stars. Also, PNe may be progenitors of Type Ia supernovae, which return prodigious amounts of iron to the ISM; a superb candidate is the pair of central stars in PN G135.9+55.9 (Tovmassian et al.\\ 2007). \\cite{miszalski09} find several similar short-period central star binaries. Thus, we need to know the number of PNe in the Galaxy in order to develop accurate models of the chemical enrichment rates since the initial burst of Type II supernovae when the Galaxy was young.  More directly, the number count of PNe in the Galaxy depends strongly on whether single stars like the Sun can form a PN, or if more unusual circumstances are required, such as the presence of a close binary companion. That is, the true number of PNe depends on how stars evolve and how PNe form, and so, a simple count provides a strong test of one of the predictions of stellar evolution -- that all low mass stars between 1--8 M$_{\\odot}$ will go through a PN phase. If this is true, there will be a much larger number of PNe in the Galaxy than if a special condition is required. The models of \\cite{moe06}, for example, predict that there are $46~000\\pm22~000$ for the general case, but only $\\sim6~600$ (De Marco \\& Moe 2005) if close binaries (e.g., common envelope phase) are required. \\cite{miszalski09} confirm earlier estimates that the binary fraction of PN central stars is only 10-20\\% of all PNe, and thus, binarity is not likely to be a major factor in the formation process. If true, there should be many PNe waiting to be found in the Galaxy.  The goal of this study is to demonstrate another method toward improving the Galactic census of PNe. We summarize the census problem in section 2, describe our technique in section 3, and report on the first results of our survey in section 4. Section 5 describes some implications of the results. ",
        "conclusions": "We have demonstrated a method for finding dozens to hundreds of faint PNe using digital sky survey data. We do not expect to find the missing thousands, or tens of thousands of PNe that are predicted to be hiding in the Galaxy, obstructed by dust. If those objects truly exist, other techniques are needed. Br$\\gamma$ is a good choice for a future emission line survey because it becomes one of the brightest lines from PNe when the visual extinction exceeds $\\sim$12 magnitudes (Jacoby \\& Van de Steene 2004). With the recent development of wide-field IR imagers, it is time to begin a survey of this type.  Still, the optical surveys continue to be useful because they probe more deeply than earlier PN surveys. An intriguing result is the emergence of a large number of very spherical PNe at low surface brightness in this and other recent surveys. Their origin and significance remains an active topic for discussion."
    },
    "0910/0910.2460_arXiv.txt": {
        "abstract": "We test the luminosity function of Milky Way satellites as a constrain for the nature of Dark Matter particles. We perform dissipationless high-resolution $N$-body simulations of the evolution of Galaxy-sized halo in the standard Cold Dark Matter (CDM) model and in four Warm Dark Matter (WDM) scenarios, with a different choice for the WDM particle mass ($m_w$). We then combine the results of the numerical simulations with semi-analytic models for galaxy formation, to infer the properties of the satellite population. Quite surprisingly we find that even WDM models with relatively low $m_{w}$ values (2-5 keV) are able to reproduce the observed abundance of ultra faint ($M_v<-9$) dwarf galaxies, as well as the observed relation between Luminosity and mass within 300 pc. Our results suggest a lower limit of $1$ keV for thermal warm dark matter, in broad agreement with previous results from other astrophysical observations like Lyman-$\\alpha$ forest and gravitational lensing. ",
        "introduction": "\\label{sec:intro}  The inflationary cold dark matter scenario gives a clear prediction for the initial fluctuation spectrum responsible for the formation of the Large Scale Structure in the Universe, with considerable power down to very small scales. As a consequence we expect the mass function of dark matter haloes to rise steeply towards low masses. N-body simulations have indeed revealed that in Cold Dark Matter (CDM) models, all Galaxy-sized haloes ($M_{DM} \\sim 10^{12} M_\\odot$) should contain a large number of embedded subhaloes that survive the collapse and virialization of the parent structure (e.g Springel \\etal 2008). Although the predicted number of substructures is in reasonable agreement with observed luminosity functions (LFs) in cluster sized haloes, in Milky Way (MW) sized haloes the number of predicted sub-haloes exceeds the number of observed satellites by at least an order of magnitude (Klypin \\etal 1999, Moore \\etal 1999). This tension between models and observations represents, together with the detailed reproduction of the star formation histories (e.g. Salvadori \\etal 2008) and mass profiles of local dwarf galaxies (e.g. Walker \\etal 2009), one of the most relevant questions for the present day theories of galaxy formation and evolution.   A possible cosmological solution to this discrepancy is to replace cold dark matter with a warm species (WDM, Bode, Ostriker \\& Turok 2001 and references therein). The warm component acts to reduce the small-scale power, resulting in fewer galactic subhaloes and lower halo central densities (Col{\\'{\\i}}n \\etal 2000,2008). One possible WDM candidate is a sterile neutrino which exhibit a significant primordial velocity distribution and thus damp primordial inhomogeneities on small scales (e.g. Hansen \\etal 2002, Abazajian \\& Koushiappas 2006 Boyarsky, Ruchayskiy, Shaposhnikov 2009). Since ground based experiments may be far from directly studying in detail the nature and properties of the actual DM particle, it is of fundamental importance to determine astrophysical constraints on the maximum free streaming length of any warm candidate.  Limits on the mass of dark matter particles can be obtained from several astrophysical observations: theoretical phase space density studies combined with observations of stellar dynamics in Milly Way satellites (e.g. Boyanovsky, de Vega, Sanchez 2008; de Vega Sanchez \\& Sanchez 2009), luminosity function of high redshift QSOs (Song \\& Lee 2009), abundance of dwarf galaxies in the Local Volume (Zavala \\etal 2009) and the size of (min)-voids around the Local Group (Tikhovov \\etal 2009). Perhaps one of the most powerful tool for constraining the matter power spectrum are Lyman-$\\alpha$ forest observations (neutral hydrogen absorption in the spectra of distant quasars, Narayanan \\etal 2000, Viel \\etal 2005). Lyman-$\\alpha$ observations allows the possibility to studying the power spectrum down to small scales and over a large range of redshifts. Using HIRES data, a lower limit of $m_{w} = 1.2$ keV has been reported by Viel \\etal (2008). Even tighter limits can be obtained combining Lyman-$\\alpha$ observations with SDSS results (see Boyarsky \\etal 2009a and references therein). These results have been recently confirmed by Miranda \\& Macci\\`o (2007) in an independent estimation of lower limits for the mass of the WDM particle based on QSO lensing observations.  In the last five years our knowledge on the number and properties of satellite galaxies around our the MW has tremendously increased thanks to the results from the Sloan Digital Sky Survey (SDSS). The homogeneous sky coverage of the SDSS provided the first determination of the volume corrected MW satellite luminosity function down to luminosities as faint as 100 $L_{\\odot}$ (Koposov \\etal 2008). Exploiting this advancement in the observational knowledge of our own Galaxy, several theoretical works have revised the problem of satellite number density, coming to the conclusion that it is possible to reconcile the observational evidences with the predictions of the standard (L)CDM scenario (e.g. Koposov \\etal 2009, Macci\\`o \\etal 2009, Li \\etal 2009a).  The aim of the present work is to expand the results presented in Macci\\`o \\etal 2010 (M10, hereafter), trying to use a combination of observational data and theoretical predictions to infer significant constraints on the allowed mass of any warm dark matter particle. We start from the findings of M10 and we re-simulate one of the DM haloes presented in that work, in a WDM scenario for different choices of $m_w$. We then combine the results with the best fit Semi Analytical Model (SAM) codes, as defined in M10, to compare the resulting prediction for the luminosity function of MW satellite in order to infer a lower limit on the WDM candidate mass. ",
        "conclusions": "\\label{sec:concl}  In this paper we compare the most recent observational data on the luminosity function of MW satellites with a suite of $N$-body simulations combined with SAM techniques, with the aim of constraining a possible ``warm'' nature for DM particle. In particular, we want to make use of the new tight observational constraints on the faint-end, which have already been proved to be of fundamental importance to test SAMs in the light of the problem of MW missing satellites (see M10)  We perform a series of $N$-body simulations with different choices for the mass of the warm dark matter particle ($m_w$), and we combine them with Semi-Analytical Models (SAMs) for galaxy formation, using the best fit solution found in M10 for the parameters describing the SAMs. We then compare the resulting statistics of MW satellites with both the predictions for a LCDM Universe and observational data.  Our results show that we are indeed able to put a lower limit on $m_w$ (1 keV), which is in agreement with previous determinations based on the galaxy power spectrum and Lyman-$\\alpha$ forest observations (Viel \\etal 2008, Boyarsky \\etal 2009a). It is worth noting that our limits are less tight with respect to the previous determinations, as a number of caveats have to be taken into account. First, only the $m_w=1$ keV realization really fails on reproducing the MW satellites luminosity function (with our ``standard'' choice of SAM parameters). Using a $m_w$ value as small as 2 keV, we still predict a luminosity function which is almost indistinguishable from the concordance LCDM cosmology. This is due to the fact that the for $m_w>1$ keV, the small-scale power modifications do not affect considerably the substructures responsible for the formation of the bulk of visible MW satellites, showing an intrinsic limit of this approach with the current observational data. Moreover, the faint-end of the LF, which provides us the strongest statistical constraints, is sensitive to the details of the description of baryonic physics included in the SAMs, and in particular to the strength of stellar feedback in low-mass halo (see M10). Therefore only models that drastically failed in reproducing the data (as $m_w=1$ keV) can be realistically ruled out. In the present work we only considered a very simple WDM model; it is worth notice that there are more complex and more physically motivated models discussed in the literature (e.g. warm+cold dark matter, Boyarsky \\etal 2009b or composite dark matter Khlopov 2006, Khlopov \\& Kouvaris 2008). These scenarios deserve special studies and can provide new, interesting hints on the nature of dark matter.  This paper is mainly a ground test on the feasibility of our approach. New data coming from the most recent multiwavelength surveys like the Panoramic Survey Telescope and Rapid Response System (Pan-STARRS)\\footnote{http://www.ps1sc.org/index.htm} will provide a better determination of the MW satellite properties and then an optimal data-set to increase our knowledge about the nature of dark matter."
    },
    "0910/0910.2710_arXiv.txt": {
        "abstract": "New low frequency radio telescopes currently being built open up the possibility of observing the 21-cm radiation before the Epoch of Reionization in the future, in particular at redshifts $200 > z > 30$, also known as the dark ages. At these high redshifts, Cosmic Microwave Background (CMB) radiation is absorbed by neutral hydrogen at its 21-cm hyperfine transition. This redshifted 21-cm signal thus carries information about the state of the early Universe and can be used to test fundamental physics. We study the constraints these observations can put on the variation of fundamental constants. We show that the 21-cm radiation is very sensitive to the variations in the fine structure constant and can in principle place constraints comparable to or better than the other astrophysical experiments ($\\Delta\\alpha/\\alpha= < 10^{-5}$). Making such observations will require radio telescopes of collecting area $10 - 10^6 \\hspace{2 pt}\\rm{km}^2$ compared to  $\\sim 1\\hspace{2 pt} \\rm{km}^2$ of current telescopes. These observations will thus provide independent constraints on $\\alpha$ at high redshifts, observations of quasars being the only alternative. More importantly the 21-cm absorption of CMB  is the only way to probe the redshift range between recombination and reionization. ",
        "introduction": "In the standard model of cosmology the electrons recombine with protons and heavy nuclei, mostly helium, to form neutral atoms at a redshift of around $z\\sim 1100$.  The mostly neutral baryonic gas then undergoes gravitational collapse to form the first stars which then reionize the Universe. The process of reionization is expected to start at around $z\\lesssim 30$ and end by $z\\sim 6$. The era between recombination and reionization has come to be known as the dark ages due to the absence of any stars or other light emitting collapsed objects. There is however the CMB which gets absorbed by the hydrogen gas at 21-cm rest wavelength corresponding to the hyperfine transition of ground state of hydrogen. The observation of these dark ages using the 21-cm line is considered to be the last frontier in cosmology due to the new information it will bring about this hitherto unexplored redshift range and also because the amount of information we can get from the 21-cm observations of dark ages is many orders of magnitude more than is possible by any other cosmological probe.  We can probe the big bang nucleosynthesis at a redshift of $z\\sim 10^9$ by measuring the abundance of light elements, this gives information only about the average properties of the Universe. CMB is a snapshot of Universe at $z\\sim 1100$ and has information on scales $\\gtrsim 10\\hspace{4 pt} \\rm{Mpc}$. Below $10\\hspace{4 pt}\\rm{Mpc}$ Silk damping causes the CMB power spectrum to drop rapidly making it hard to measure. Observations of galaxies, clusters of galaxies and $\\rm{Ly}\\alpha$ forest probes $z\\sim >6$ also on scales $> 10\\hspace{4 pt}\\rm{Mpc}$. On small scales the matter perturbations at low redshifts become non-linear making it hard to reconstruct the  initial conditions or constrain fundamental physics. The 21-cm observations will allow us to probe a new redshift range of $30 \\gtrsim z \\gtrsim 200$, which is not accessible by any other method,   on scales down to $\\sim \\rm{few}\\hspace{4 pt} \\rm{kpc}$ and use that information to constrain initial conditions and fundamental physics. ",
        "conclusions": "We have shown that the 21-cm observations of the dark ages can provide very tight constraints on the variation of the fine structure constant. In addition we should also expect similar sensitivity to the electron to proton mass \\begin{figure} \\includegraphics[clip=true]{sait_khatri_FIG3.eps} \\caption{ \\footnotesize Constraints on $\\alpha$ from 21-cm observations of dark ages for different radio telescopes labeled by collecting area. Also shown is a current telescope, LOFAR, that is currently operating in Europe. } \\label{fish} \\end{figure} ratio. This is easily seen by observing that $A_{10}\\propto \\frac{1}{m_e^2}\\left(\\frac{m_em_p}{m_e+m_p}\\right)^3$, $\\sigma_T\\propto \\frac{1}{m_e^2}$ etc. The spin change collision cross section also depend on $m_e/m_p$ through the kinetic term in the electronic Hamiltonian of the hydrogen molecule. One of the challenges in observing this 21-cm cosmological signal is the presence of the synchrotron foregrounds. The sky noise due to the synchrotron foregrounds in the galaxy is expected to be $T_{sky}\\sim 19000K\\left(\\frac{22MHz}{\\nu}\\right)^{2.5}$ \\cite{rog}. This is many orders of magnitude larger than the cosmological signal which is of the order of $\\sim  mk$. The foregrounds can still be removed from the cosmological signal due to the following reason \\cite{zal}. The foregrounds are expected to be correlated in frequency while the cosmological signal traces a gaussian random field and would be uncorrelated for frequencies sufficiently separated. Thus the detection of the cosmological signal is challenging but possible. Interstellar scattering may limit the smallest scales that are observable to $\\ell \\lesssim 10^6$ \\cite{cohen}. Terrestrial EM interference from radio/TV/communication and  Earth's ionosphere poses problems for telescopes on ground. Earth related problems may be solved by going to the Moon and there are proposals for doing so, one of which is the Dark Ages Lunar Interferometer (DALI) \\cite{dali}. In conclusion 21-cm cosmology promises a large wealth of data with which to constrain and learn about fundamental physics."
    },
    "0910/0910.0715_arXiv.txt": {
        "abstract": "We study observational implications of the stochastic gravitational wave background and a non-Gaussian feature of scalar perturbations  on the curvaton mechanism of the generation of density/curvature fluctuations, and show that they can determine the properties of the curvaton in a complementary manner to each other.  Therefore even if Planck could not detect any non-Gaussianity, future space-based laser interferometers such as DECIGO or BBO could practically exhaust its parameter space. ",
        "introduction": "The idea of inflation \\cite{lindebook} has become a standard paradigm since it was proposed in the early 1980's. Inflation is an accelerated expansion epoch in the very early Universe, which solves the flatness, the horizon and also the monopole problems naturally. Furthermore quantum fluctuation of the inflaton $\\phi$, which is a scalar field responsible for the accelerated expansion, can provide the seed of cosmic density/curvature fluctuations \\cite{yuragi} observed through the cosmic microwave background (CMB) anisotropy  \\cite{Komatsu:2008hk}, galaxy clustering, etc.   Another prediction of inflation, which is in fact more generic, is the generation of  the stochastic gravitational wave background or the tensor perturbation, whose spectrum is nearly scale-invariant \\cite{staro,Allen:1987bk,Turner:1990rc,Sahni:1990tx,Turner:1993vb,Turner:1993xz,Liddle:1993zj,Turner:1995ge,Turner:1996ck,Allen:1997ad,Seto:2003kc, Weinberg:2003ur,Tashiro:2003qp,Ungarelli:2005qb,Smith:2005mm,Seto:2005qy,Kudoh:2005as,Boyle:2005se,Smith:2006xf,Chongchitnan:2006pe,Friedman:2006zt,Zhao:2006is,Watanabe:2006qe,Chiba:2007kz,Boyle:2007zx,jyessay,Smith:2008pf,Nakayama:2008ip,Nakayama:2008wy, Kuroyanagi:2008ye,Mangilli:2008bw} (see Ref.~\\cite{Maggiore:1999vm} for a review). The amplitude of the gravitational wave is simply proportional to the Hubble parameter, $H_{\\rm inf}$, during inflation and many inflation models predict detectable amplitude of gravitational waves. In particular, future space-based gravitational wave detectors such as DECIGO \\cite{Seto:2001qf} and/or BBO have a chance to detect inflationary gravitational wave background. Recently, it was pointed out that thermal history of the early universe is imprinted in the spectral shape of the gravitational wave background \\cite{Seto:2003kc,Boyle:2005se,Boyle:2007zx,jyessay, Nakayama:2008ip,Nakayama:2008wy}. In particular, the reheating temperature of the Universe after inflation can be determined or constrained from future gravitational wave experiments \\cite{jyessay,Nakayama:2008ip,Nakayama:2008wy}. Thus any detection of primordial gravitational wave background gives useful information on the early Universe.   Traditionally, the spectrum of the comoving curvature perturbation is parametrized as \\begin{equation} \\Delta^2_{\\zzeta}(k)=\\Delta^2_{\\zzeta}(k_*)\\left ( \\frac{k}{k_*} \\right )^{n_s-1}. \\end{equation} Here $k_\\ast$ is the pivot scale which we take $k_*=0.002$~Mpc$^{-1}$  and $n_s$ is the scalar spectral index. Inflation predicts nearly scale-invariant spectrum with $n_s\\cong 1$, with its precise value determined by the shape of the potential \\cite{yuragi}, and it agrees well with observations so far~\\cite{Komatsu:2008hk}. If $\\phi$ has a single component with a canonical kinetic term, the resultant curvature fluctuation is Gaussian distributed, which is again in agreement with the observation today \\cite{Komatsu:2008hk}.  However, the origin of the density perturbation is not limited to the quantum fluctuation of the inflaton. Another scalar field, called curvaton, may be responsible for the generation of the observed density perturbation \\cite{Mollerach:1989hu,Lyth:2001nq}. It is undoubtedly an important task to distinguish them in order to understand the physics of the early Universe. One such signature may come from the non-Gaussianity in the CMB anisotropy \\cite{Bartolo:2004if}, since the curvaton scenario can produce large enough non-Gaussian feature to be detected \\cite{Lyth:2002my,Enqvist:2005pg}, while standard inflation models predict negligible non-Gaussianity \\cite{Acquaviva:2002ud,Maldacena:2002vr,Seery:2005wm,Yokoyama:2007uu}. However, the situation where the curvaton generates large non-Gaussianity is somewhat limited and it is possible that the curvaton accounts for the observed density perturbation without generating large non-Gaussianity.  In this paper, we point out that the observation of stochastic gravitational wave background plays a very important role to probe the physics of the curvaton scenario. In the case that the curvaton generates negligible non-Gaussianity, there must be an entropy production process by the curvaton decay itself, and such a non-standard thermal history is imprinted in the spectrum of the inflationary gravitational wave background. Future space-based gravitational wave detectors, such as DECIGO and/or BBO may be able to do this job. This opens up a possibility to find an evidence of the curvaton scenario.  Thus measurement of the non-linearity parameter $f_{\\rm NL}$ which is a simplified measure of non-Gaussianity and that of tensor perturbations by DECIGO/BBO play complementary roles to each other.  The rest of the paper is organized as follows. In Sec.~\\ref{sec:curvaton}, the gravitational wave background spectrum and non-Gaussianity in the curvaton scenario are summarized. In Sec.~\\ref{sec:probe} the detection possibility is discussed. Sec.~\\ref{sec:conc} is devoted to the conclusion. ",
        "conclusions": "\\label{sec:conc}   In this paper we have investigated a possible observable signatures from curvaton scenarios, including non-Gaussianity, detection of tensor modes in CMB and direct detection of gravitational wave background. First the measurement of the tensor power spectrum by CMB is very important to identify the curvaton scenario through the consistency relation (\\ref{consistency}). If the curvaton once dominates the Universe, an entropy production process by its decay is imprinted in the gravitational wave background spectrum and can be confirmed by future space-based laser interferometer experiments. For the opposite case where the curvaton is a subdominant component before it decays, a large non-Gaussianity is predicted, which may be confirmed by Planck experiment. In this sense, the gravitational wave signal and non-Gaussian signal are complementary to each other. This allows us to probe the curvaton scenario for large parameter spaces, giving information on the properties of the curvaton such as its decay rate and amplitude, etc. It will in turn provide important information on high-energy physics beyond the reach of terrestrial experiments."
    },
    "0910/0910.0838_arXiv.txt": {
        "abstract": "We study the Spitzer Infrared Array Camera (IRAC) mid-infrared (rest-frame optical) fluxes of 14 newly WFC3/IR-detected $z\\sim7$ $z_{850}-$dropout galaxies and 5 $z\\sim8$ $Y_{105}$-dropout galaxies. The WFC3/IR depth and spatial resolution allow accurate removal of contaminating foreground light, enabling reliable flux measurements at $3.6\\mu$m and $4.5\\mu$m. None of the galaxies are detected to $[3.6]\\approx26.9$ (AB,~$2\\sigma$), but a stacking analysis reveals a robust detection for the $z_{850}$-dropouts and an upper limit for the $Y_{105}-$dropouts. We construct average broadband spectral energy distributions using the stacked ACS, WFC3, and IRAC fluxes and fit stellar  population synthesis models to derive mean redshifts, stellar masses, and ages. For the $z_{850}-$dropouts, we find $z=6.9^{+0.1}_{-0.1}$, $(U-V)_{rest}\\approx0.4$, reddening $A_V=0$, stellar mass $<M^*>=1.2^{+0.3}_{-0.6}\\times10^9~M_\\sun$ (Salpeter initial mass function). The best-fit ages $\\sim300~$Myr, $M/L_V\\approx0.2$, and $SSFR\\sim1.7$~Gyr$^{-1}$ are similar to values reported for luminous $z\\sim7$ galaxies, indicating the galaxies are smaller but not much younger. The sub$-L^*$ galaxies observed here contribute significantly to the stellar mass density and under favorable conditions may have provided enough photons for sustained reionization at $7<z<11$. In contrast, the $z=8.3^{+0.1}_{-0.2}$  $Y_{105}$-dropouts have stellar masses that are uncertain by 1.5 dex due to the near-complete reliance on far-UV data.  Adopting the $2\\sigma$ upper limit on the $M/L(z=8)$, the stellar mass density to $M_{UV,AB}<-18$ declines from $\\rho^*(z=7)=3.7^{+1.4}_{-1.8}\\times10^6~M_\\sun~$Mpc$^{-3}$ to $\\rho^*(z=8)<8\\times10^5~M_\\sun~$Mpc$^{-3}$, following $\\propto(1+z)^{-6}$ over $3<z<8$. Lower masses at $z=8$ would signify more dramatic evolution, which can be established with deeper IRAC observations, long before the arrival of the James Webb Space Telescope. ",
        "introduction": "The reionization epoch represents the latest observational frontier of galaxy formation. Little is known about the galaxy properties during this time, including feedback, metal production, and their contribution to reionization. Extensive studies have been done of galaxies selected by the dropout technique out to $z=6$ (see e.g., \\citealt{St03, Ya04, Bo06, Mc09a}). Pushing these studies to $z\\gtrsim7$ has proven extraordinarily challenging: only $\\sim$25 high-quality candidates are currently known \\citep[][R. J. Bouwens et al. in preparation]{Bo08,Oe09a,Ca09,Ou09}.  The arrival of WFC3/IR aboard {\\it Hubble Space Telescope (HST)} heralds a dramatic improvement in our ability to survey the reionization era, allowing us to identify the predominant low-luminosity galaxies at $z\\gtrsim7$ by their redshifted $UV$ light. Whereas {\\it HST} remains crucial for identifying the galaxies, longer-wavelength data are necessary for constraining the stellar masses and ages. Access to the rest-frame optical is offered by the Infrared Array Camera (IRAC; \\citealt{Fa04}) on {\\it Spitzer}, which by itself has permitted estimates of stellar masses for large numbers of $z\\sim5-6$ sources \\citep{Ey05, Ya06, St09}, and even a few at $z\\sim7$ \\citep{Eg05,La06}. One surprising early finding was the number of quite massive $\\sim10^{10}M_\\sun$ galaxies with appreciable ages ($200-300$Myr) suggesting the galaxies formed substantial amounts of stars at even earlier times, well into the epoch of reionization \\citep{St07, Ya06}.  \\figstack  In this Letter, we extend the stellar mass and age estimates to lower luminosities ($0.06~L^*_{z=3}$) and higher redshift $z\\approx8$ by analyzing the IRAC mid-IR fluxes of the robust $z_{850}$-dropout and $Y_{105}$-dropout galaxy samples found by \\citet{Oe09b} and \\citet{Bo09a} in the newly acquired WFC3/IR data  (GO-11563, PI Illingworth) over the Hubble Ultra Deep Field (HUDF; see also \\citealt{Bu09,Mc09b}). We adopt an $\\Omega_M=0.3, \\Omega_\\Lambda=0.7,$ cosmology with $H_0=70$~km~s$^{-1}$Mpc$^{-1}$. Magnitudes are in the AB photometric system \\citep{Ok83}. Imaging depths refer to $1\\sigma$ AB point source total magnitude, unless stated otherwise. ",
        "conclusions": "The reionization epoch represents the latest frontier of galaxy formation theory, which is now being probed to extreme depths with the amazingly efficient WFC3/IR camera aboard {\\it HST}. Using the conservative $z_{850}-$dropout and \\ydrop \\ samples of \\citet{Oe09b} and \\citet{Bo09a} we find {\\it Spitzer}/IRAC capable of an equally remarkable feat: the stacked $3.6\\micron$ detection of even the faintest $z\\approx6.9$ galaxies found by WFC3/IR, providing direct proof that IRAC can probe much deeper than widely accepted. In addition we derive upper limits for the newly discovered $z\\approx8.3$ galaxies, some of which may well be at higher redshift than the currently most distant $z\\approx8.2\\pm0.1$ GRB \\citep{Sa09,Ta09}. We now elaborate on various implications.\\\\  {\\it 1) Stellar populations at $z\\sim7$} \\\\ {\\it Spitzer}/IRAC places two important constraints on the stellar populations. First, the combination of the blue far-UV slope and the moderately red $H_{160}-[3.6]=0.7^{+0.2}_{-0.25}$ color in the ultrafaint ($H_{160}\\approx27.9$)  $z\\sim7$ galaxies indicates the presence of a modest Balmer break, expected for evolved  stellar populations. Very young stellar ages are ruled out (CSF age$_w\\gtrsim80~$Myr at $95$\\%). Interestingly, low-metallicity models, which produce bluer far-UV colors at fixed age, provide significantly better fits than do Solar, but for both it is challenging to reproduce the extremely blue observed far-UV slope ($\\beta=-3.0\\pm0.2$ as found by \\citealt{Bo09b}). Second, the average stellar masses $<M> \\sim 1\\times10^{9}~M_\\sun$ indicate appreciable mass-to-light ratios M/L$_{1500}\\approx0.1$ and M/L$_V\\approx0.2$. The ratios are comparable those of the more massive $\\sim7\\times 10^9~M_\\sun$ sample of \\cite{Go09}, suggesting that the $M/L$ ratios do not strongly depend on luminosity or mass. We find no appreciable difference between the best-fit parameters of BC03, MA05, and CB07 models, although IMF variations are a potentially significant source of systematic uncertainty: a steeper high-mass slope (e.g., \\citealt{We05}) would lower the inferred stellar ages, a higher characteristic mass (e.g., ``bottom light'', \\citealt{vD08}) would lower the total stellar masses and $M/L$s, and any evolution with time would complicate comparisons between SFR and stellar mass across epochs.\\\\  {\\it 2) Star formation histories} \\\\ SFHs are notoriously hard to constrain from observations of individual galaxies. However, statistical constraints can be obtained by comparing the SSFRs as a function of time and mass, as different shapes of the SFH give rise to distinctly different evolution (e.g., \\citealt{La07}). One remarkable recent result is the observation that the SSFR at fixed $\\sim5\\times10^9 M_\\sun$ does not evolve from $z=3-7$ \\citep{St09,Go09}. Here we benefit from the extreme depths of the WFC3/IR data to probe the lowest luminosities, suggesting that the SSFR at $z\\sim7$ is also not strongly dependent on luminosity or stellar mass over the range $1-7\\times10^9 M_\\sun$. This may imply that the SFR correlates with stellar mass with a logarithmic slope close to 1. Both results (with redshift and mass) qualitatively agree with numerical simulations for early galactic star formation ($z>4)$, which robustly predict that SFR is proportional to M*, with similar SFHs in haloes of different masses \\citep{Da08}, and that the SFRs of individual galaxies are constantly rising \\citep{Fi07}.\\\\  {\\it 3) Evolution of the stellar mass density to $z=7-8$} \\\\ The low-luminosity galaxies probed in this paper are expected to contribute significantly to the stellar mass density owing to their substantial $M/L$. Following the approach of \\citet{Go09}, we derive integrated stellar mass densities at $z=7-8$ by multiplying the UV-luminosity densities of \\citet{Bo09a}, integrated to $M_{UV,AB}=-18$, with the mean $M/L$ derived for the galaxy samples in this Letter. For the \\zdrop s at $z\\sim7$, this results in $\\rho^*=3.7^{+1.4}_{-1.8}\\times10^6~M_\\sun $Mpc$^{-3}$. The evolution of the mass density is shown in Figure~\\ref{figmass}, compared to previous luminous samples, which were corrected to the same luminosity limit assuming a constant $M/L$. The stellar masses at $z=8$ are uncertain by $\\sim1.5~$dex due to the weak constraints from IRAC, so instead of plotting the best-fit $M/L$ we show the  95\\% confidence upper limit. The integrated mass density  to a limit of $M_{UV,AB}<-18$ then declines to $<8\\times10^5~M_\\sun$Mpc$^{-3}$. This is marginally consistent with the extrapolation of the $\\propto(1+z)^{-6}$ evolution over the range  $3<z<7$. The upper limit may be evidence for a more rapid evolution toward $z=8$, but deeper IRAC data and larger samples are needed to establish this firmly. The uncertainties quoted above do not include cosmic variance, which is estimated to be $\\sim30$\\% for the $z\\sim7$ sample and $\\sim40$\\% at $z\\sim8$ \\citep{Tr08}.\\\\  {\\it 4) Reinionization since $z\\sim11$} \\\\ Combined, the results have profound implications for the ability of the universe to be reionized by photons from star formation -- and to remain so for an extended period. Crucially, the large stellar mass density in $z\\sim7$ low-luminosity galaxies points to significant amounts of star formation at earlier times. We can therefore ask whether this corresponds to enough ionizing photons at higher redshift to keep the universe reionized. Let us assume the stellar mass density $\\rho^*=3.7\\times10^6~M_\\sun$Mpc$^{-3}$ was assembled in the $340$~Myr between $7<z<11$, in the reionization era (WMAP5, \\citealt{Ko09}). Correcting by 10\\% to account for mass loss in stellar evolution, the possible sustained SFR density would then be $\\rho_{SFR}\\approx0.012~M_\\sun$yr$^{-1}$Mpc$^{-3}$, significantly higher than that has been derived for luminous galaxies alone \\citep{Go09}. However, the primary uncertainty is the fiducial critical SFR density needed for reionization $\\rho^{SFR}_{crit}\\propto~C/f_{esc}$ \\citep{Ma99}, which depends sensitively on the values for the \\ion{H}{2} IGM clumping factor $1<C<30$ and the fraction of ionizing photons $0.05<f_{esc}<1$ that leak unhindered out of the galaxies (e.g., \\citealt{Ou09}). The popular choice $(C=30,f_{esc}=0.1)$ leads to $<\\rho^{SFR}_{crit}>=0.7~M_\\sun$yr$^{-1}$Mpc$^{-3}$ averaged over $7<z<11$, which is $60$ times higher than that can be explained by the stellar masses. However, at early times the clumping factor may be as low as $C=3-6$ (e.g., \\citealt{Tr07,Bo07,Pa09}). The escape fraction may be significantly higher as well: recent numerical simulations suggest $f_{esc}\\gtrsim0.2$ \\citep{Wi09,Pa09,Ya09} and \\citet{Bo09b} note that the extremely blue far-UV slope of the galaxies studied here could also result from large escape fractions $f_{esc}\\gtrsim0.3$. The choice $(C=3,f_{esc}=0.5)$ leads to $<\\rho^{SFR}_{crit}>\\approx0.014~M_\\sun$yr$^{-1}$Mpc$^{-3}$(over $7<z<11$), which means photons from the low-luminosity star forming galaxies observed here are capable of causing sustained reionization since $z=11$. \\\\  The first {\\it Spitzer}/IRAC detection of low-luminosity $z=7$ galaxies and their red $H-[3.6]\\approx0.7$ colors provide important insights into the earliest phases of galaxy evolution, showing that these early galaxies have relatively high $\\sim300~$Myr ages and $M/L_V\\approx0.2$. Consequently, low-luminosity galaxies contribute significantly to the stellar mass density $\\rho^*(z=7)=3.7^{+1.4}_{-1.8}\\times10^6~M_\\sun $Mpc$^{-3}$, suggesting that the SFR corresponding to that mass may have provided a substantial fraction of the ionizing photons that kept the universe reionized at earlier times. Future IRAC studies of larger numbers of high redshift galaxies are required to bolster the results in this paper and to further constrain the evolution of the mass density toward the highest redshifts $z\\gtrsim8$. The {\\it Spitzer} post-cryogenic phase (``warm mission'') provides a rare window of opportunity to observe these galaxies in a timely manner to the $\\sim0.5~$mag deeper limits needed, long before the arrival of the {\\it James Webb Space Telescope}.\\\\"
    },
    "0910/0910.5901_arXiv.txt": {
        "abstract": "We have used archival {\\it HST} H$\\alpha$ images to study the immediate environments of massive and intermediate-mass young stellar object (YSO) candidates in the Large Magellanic Cloud (LMC).  The sample of YSO candidates, taken from \\citet{gruendl08}, was selected based on {\\it {\\it Spitzer}} IRAC and MIPS observations of the entire LMC and complementary ground-based optical and near-infrared observations. We found {\\it HST} H$\\alpha$ images for 99 YSO candidates in the LMC, of which 82 appear to be genuine YSOs.  More than 95\\% of the YSOs are found to be associated with molecular clouds.  YSOs are seen in three different kinds of environments in the H$\\alpha$ images: in dark clouds, inside or on the tip of bright-rimmed dust pillars, and in small \\ion{H}{2} regions.  Comparisons of spectral energy distributions for YSOs in these three different kinds of environments suggest that YSOs in dark clouds are the youngest, YSOs with small \\ion{H}{2} regions are the most evolved, and YSOs in bright-rimmed dust pillars span a range of intermediate evolutionary stages. This rough evolutionary sequence is substantiated by the presence of silicate absorption features in the {\\it Spitzer} IRS spectra of some YSOs in dark clouds and in bright-rimmed dust pillars, but not those of YSOs in small \\ion{H}{2} regions. We present a discussion on triggered star formation for YSOs in bright-rimmed dust pillars or in dark clouds adjacent to \\ion{H}{2} regions. As many as 50\\% of the YSOs are resolved into multiple sources in high-resolution {\\it HST} images.  This illustrates the importance of using high-resolution images to probe the true nature and physical properties of YSOs in the LMC. ",
        "introduction": "The {\\it Spitzer Space Telescope}, with its high angular resolution and sensitivity at mid-infrared wavelengths, has made it possible for the first time to survey massive young stellar objects (YSOs) in nearby external galaxies.  In particular, in the Large Magellanic Cloud (LMC), because of its close proximity (50 kpc; \\citealt{feast99}) and low inclination ($\\sim$ 30$\\arcdeg$; \\citealt{nikolaev04}), massive and intermediate-mass YSOs can be resolved by {\\it Spitzer} and inventoried throughout the entire galaxy.  The LMC was observed by {\\it Spitzer} under a Legacy Program, Surveying the Agents of Galaxy Evolution (SAGE), that mapped the central 7\\arcdeg\\ $\\times$ 7\\arcdeg\\ area of the galaxy \\citep{meixner06}. \\citet{gruendl08} made use of the archival {\\it Spitzer} data from the SAGE survey along with complementary ground-based optical and near-infrared observations and identified a comprehensive sample of massive and intermediate-mass YSO candidates of the entire LMC. This sample consists of a total 855 {\\it definite}, 317 {\\it probable}, and 213 {\\it possible} massive ($>$ 10 $M_\\sun$) and intermediate-mass ($>$ 4 $M_\\sun$) YSO candidates.  Most of these YSO candidates are found to be concentrated in or around molecular clouds or \\ion{H}{2} regions \\citep{gruendl08}.  While {\\it Spitzer} enables the detection of individual massive and intermediate-mass YSOs in the LMC \\citep[e.g.,][]{chu05,jones05,caulet08,whitney08,gruendl08}, the angular resolution of {\\it Spitzer} is not sufficient to resolve multiple sources within 1--2$\\arcsec$ or allow a close look at the environments of these massive YSOs. The {\\it Hubble Space Telescope (HST)} on the other hand, with an angular resolution ten times higher than that of {\\it Spitzer}, reveals sub-arcsec size features and hence is very useful for a detailed examination of the environments of these massive YSOs.  We have therefore searched through the {\\it HST} archive for continuum and H$\\alpha$ images that cover LMC YSO candidates from the above catalog. In many cases, the exposure times of the continuum images are short and hence not very useful, as noted by \\citet{chen08}.  The H$\\alpha$ images have longer exposure times and are very useful in revealing interstellar environments of YSOs in \\ion{H}{2} complexes where the interstellar gas is photoionized by antecedent massive stars.  For example, ionization fronts on the surface of dense molecular clouds can be recognized by the sharply enhanced H$\\alpha$ surface brightness.   H$\\alpha$ images are also useful in revealing circumstellar environments of massive YSOs.  For example, the UV flux of a YSO may photoionize its circumstellar medium and form a compact \\ion{H}{2} region, or outflows from a YSO may produce observable features like Herbig-Haro objects \\citep{heydari99,chu05}.  Thus, in this paper we make use of primarily the {\\it HST} H$\\alpha$ images to examine the immediate surroundings of the LMC YSO candidates.  The continuum images are only used to supplement the analyses of stellar properties for a few YSOs.  The remaining paper is organized as follows:  Section 2 describes the datasets used in this work and our method of analysis, Section 3 describes the environments of YSO candidates, Section 4 discusses the properties of small \\ion{H}{2} regions discovered around some massive YSO candidates, Section 5 describes the mid-infrared spectral characteristics of some YSO candidates, Section 6 discusses the spectral energy distributions (SEDs) of the YSO candidates, Section 7 presents a discussion on triggered star formation, and finally Section 8 summarizes the paper. ",
        "conclusions": "We have used archival {\\it HST} H$\\alpha$ images to examine the immediate environments of massive and intermediate-mass YSO candidates of the LMC. In total, archival {\\it HST} H$\\alpha$ images were found for 99 YSO candidates. By examining the source morphology and the environment of YSO candidates in multiwavelength images, we conclude that 82 of the candidates are genuine YSOs. All of these 82 YSOs are found in only three different environments, in dark clouds, in bright rimmed dust pillars, or in small \\ion{H}{2} regions. Forty-nine YSOs are in dark clouds, 34 of which show no optical counterparts, and the remaining 15 show faint optical counterparts in {\\it HST} H$\\alpha$ images. Nineteen YSOs are associated with bright-rimmed dust pillars, of which five are in dark clouds, and three show small \\ion{H}{2} regions. Twenty two YSOs show small \\ion{H}{2} regions, of which 8 are well resolved, 4 are marginally resolved, and the remaining 10 are unresolved in the {\\it HST} H$\\alpha$ images.  We calculate observed (reddened) H$\\alpha$ fluxes of the resolved \\ion{H}{2} regions and present the estimated spectral types of these massive YSOs.  For the marginally resolved as well as the unresolved \\ion{H}{2} regions, we use $UBV$ photometry or {\\it HST} continuum images to estimate the stellar continuum fluxes in the H$\\alpha$ passband, and demonstrate that most of them possess H$\\alpha$ line emission, indicating that they indeed have small \\ion{H}{2} regions.  YSOs for which {\\it Spitzer} IRS spectra are available, predominantly show PAH emission and/or fine-structure lines in the spectra. Nine YSOs associated with dark clouds, with two of them also in bright-rimmed dust pillars, show silicate absorption features in the IRS spectra, whereas none of the YSOs associated with small \\ion{H}{2} regions show silicate absorption features. YSOs in small \\ion{H}{2} regions including the marginally resolved, and the unresolved ones, show PAH emission and/or fine-structure lines in the IRS spectra.  The comparison of YSO environments with the SEDs reveal an evolutionary sequence of YSOs in different environments.  YSOs in dark clouds with no optical counterparts are mostly Type I and hence the youngest.  YSOs in dark clouds but with optical counterparts are mostly Type II and more evolved compared to YSOs with no optical counterparts.  YSOs in bright-rimmed dust pillars are either Type II or Type II/III and are in the intermediate stage. YSOs with resolved \\ion{H}{2} regions are mostly Type III and are the most evolved.  Finally, YSOs with marginally resolved \\ion{H}{2} regions, and unresolved \\ion{H}{2} regions are a mix of Type II and Type III, and should be on average younger compared to YSOs with resolved \\ion{H}{2} regions.  As many as 50\\% YSOs are resolved into multiple sources when seen in {\\it HST} images, signifying the importance of using high-resolution images to probe the true nature of YSOs and to study their immediate environments.  We investigate the issue of triggered star formation for YSOs in bright-rimmed dust pillars and YSOs in dark clouds adjacent to \\ion{H}{2} regions. The thermal pressures of ionized surfaces of bright-rimmed dust pillars are found to be much higher compared to typical thermal pressures of dust globules.  Thus the YSOs in bright-rimmed dust pillars have likely formed due to triggering from external pressure. Finally, we show that by examining the immediate environments of the YSOs using the high-resolution {\\it HST} images, we can learn about the evolutionary stages of massive YSOs.  A systematic survey of massive YSOs using {\\it HST} will be very useful to study the evolutionary aspects of massive YSOs and to understand star formation in wide range of interstellar environments."
    },
    "0910/0910.0009_arXiv.txt": {
        "abstract": "We present a new method to directly measure the opacity from \\ion{H}{1} Lyman limit (LL) absorption \\kll\\ along quasar sightlines by the intergalactic medium (IGM). The approach analyzes the average (``stacked'') spectrum of an ensemble of quasars at a common redshift to infer the mean free path \\mfp\\ to ionizing radiation.  We apply this technique to 1800 quasars at $z = 3.50 - 4.34$ drawn from the Sloan Digital Sky Survey (SDSS), giving the most precise measurements on \\kll\\ at any redshift. From $z=3.6$ to 4.3, the opacity increases steadily as expected and is well parameterized by $\\mmfp = \\bestmfp - \\slopemfp (z-3.6)$ with $\\bestmfp = (\\vbestmfp) \\, \\umfp$ and $\\slopemfp = (\\vslopemfp) \\, \\umfp$ (proper distance). The relatively high \\mfp\\ values indicate that the incidence of systems which dominate \\kll\\ evolves less strongly at $z>3$ than that of the \\lya\\ forest. We infer a mean free path three times higher than some previous estimates, a result which has important implications for the photo-ionization rate derived from the emissivity of star forming galaxies and quasars. Finally, our analysis reveals a previously unreported, systematic bias in the SDSS quasar sample related to the survey's color targeting criteria. This bias potentially affects all $z \\sim 3$ IGM studies using the SDSS database. ",
        "introduction": "The observed high transmission of $z \\sim 3$ quasars at rest wavelengths \\lbr\\ blueward of \\ion{H}{1} \\lya\\ reveals that the intergalactic medium (IGM) is highly ionized \\citep{gp65}. The presence of the \\lya\\ forest demands an intense, extragalactic ultraviolet background (EUVB) radiation field. The quasars themselves provide a significant fraction of the required ionizing flux, buoyed by the emission from more numerous yet fainter star-forming galaxies. Several recent studies have argued that the latter population dominates the EUVB at $z \\gtrsim 3$ \\citep{flh+08,cbt09,dww09}, where the quasar population likely declines \\citep[e.g.][]{fan04}. These assertions, however, hinge on the opacity of the IGM to ionizing radiation via \\ion{H}{1} Lyman limit absorption (\\kll) which directly impacts estimates of the EUVB measured from the integrated quasar and stellar ionizing emissivity.  Traditionally, \\kll\\ has been estimated from the incidence of so-called Lyman limit systems (LLSs) via surveys of quasar spectroscopy \\citep[e.g.][]{lzt91,storrie94,peroux_dla03}. Observationally, one can rather easily identify systems with large optical depths $\\mtll \\gtrsim 2$ and the majority of these surveys have probed to this limit. The integrated opacity of the IGM, however, includes and is likely dominated by gas with $\\mtll < 1$, the so-called partial Lyman limit systems and \\lya\\ forest clouds. Systems with these \\ion{H}{1} column densities ($\\mnhi \\approx 10^{14} - 10^{17} \\cm{-2}$) are especially difficult to survey because the strong lines of the Lyman series (e.g.\\ \\lya, \\lyb) lie on the flat portion of the curve-of-growth and they exhibit only weak absorption at the Lyman limit. Therefore, current estimates of \\kll\\ are based on an extrapolation/interpolation of the frequency of systems with $\\mnhi < 10^{14} \\cm{-2}$ and $\\mnhi > 10^{17.5} \\cm{-2}$ \\citep{mhr99,sb03,flz+09}. Current constraints on \\kll\\ span over a magnitude of uncertainty, especially at $z \\gtrsim 4$.   In this Letter, we introduce a new technique to estimate \\kll\\ that avoids the traditional line-counting statistics of the IGM. We analyze the average rest-frame spectra of 1800 $z>3.5$ quasars drawn from the Sloan Digital Sky Survey, Data Release 7 \\citep{sdssdr7}. As an ensemble at a common redshift, these ``stacked'' spectra show the exponential drop in flux at $\\lambda < \\mlll = c/\\mnull = 911.76$\\AA\\ from the integrated opacity of the IGM.  We precisely evaluate \\kll\\ in a series of small redshift intervals covering $z \\approx 3.6 - 4.3$ to explore redshift evolution. We adopt a cosmology with $H_0 = 72 \\, h_{72} \\, \\mkms \\, \\rm Mpc^{-1}$, $\\Omega_{\\rm m} = 0.3$, and $\\Omega_\\Lambda = 0.7$ and report proper lengths unless specified.  \\begin{figure*} \\begin{center} \\includegraphics[height=6.8in,angle=90]{f1.eps} \\end{center} \\caption{The stacked spectrum (normalized at $\\mlbr = 1450$\\AA) from 150 mock quasar spectra with IGM absorption derived from an assumed \\fnhi\\ distribution. The adopted emission redshifts and S/N of the spectra were taken to match our SDSS quasar sample at $z=[3.59,3.63]$. The sharp decline in the flux at $\\mlbr \\approx 912$\\AA\\ is associated first with the higher order Lyman series opacity (for $\\mlbr \\gtrsim 912$\\AA) and then Lyman limit continuum opacity (for $\\mlbr < 912$\\AA). The dashed (cyan) curve overplotted on the spectrum is the predicted relative flux corresponding to \\teff\\ (Equation~\\ref{eqn:teff}), normalized at $\\mlbr = 912$\\AA. On top of this curve is the best-fit model of the flux evaluated over the interval $\\mlbr = 848-905$\\AA\\ using our new methodology (Equation~\\ref{eqn:newteff}). The resulting estimates for the mean free path \\mfp\\ are indicated on the figure and are in excellent agreement. Finally, the dotted curves show the predicted models after offsetting \\mfp\\ by $\\pm 10 \\, \\rm Mpc$ and indicate the sensitivity of our technique to measuring \\mfp. } \\label{fig:mock} \\end{figure*} ",
        "conclusions": "\\label{sec:discuss}  Our new technique provides the most precise measurements on the IGM opacity to \\ion{H}{1} ionizing radiation at any redshift. Previous estimates of \\kll\\ were limited by large uncertainties in the \\nhi\\ frequency distribution, especially the incidence of systems with $\\mnhi \\approx 10^{17} \\cm{-2}$ \\citep{mhr99,sb03,me03,mw04,flh+08}.  Our results indicate that systems dominating \\kll\\ at $z>3.6$ must evolve more slowly than the $\\mnhi \\le 10^{14} \\cm{-2}$ \\lya\\ forest. In turn, we infer a flattening in the \\ion{H}{1} frequency distribution between $\\mnhi \\approx 10^{15}-10^{17} \\cm{-2}$ \\citep{petit93,opb+07}. We provide a full discussion on the implications for \\fnhi\\ in \\cite{pow09}.  Our analysis also gives the first direct description of the evolution in \\mfp, albeit over a small redshift interval.  Our results are well parameterized by a linear decrease in \\mfp\\ but can also be described by a $(1+z)^{-\\gamma}$ power-law with $\\gamma = 3.5 - 5.5$.  These values are consistent with the observed evolution in the incidence of LLS \\citep{pow09}.  Our results refine recent inferences \\citep[e.g.][]{flh+08,dww09} that galaxies contribute significantly to the EUVB at $z>3$. We assume the photoionization rate at $z=4$ inferred from the effective \\lya\\ opacity of the IGM \\citep[$\\log\\mgigm = -12.3$;][]{fpl+08}.  Comparing this value against the photoionization rate inferred from quasars \\gq\\ using an emissivity $\\meq = 2\\sci{24} \\; {\\rm erg s^{-1} Hz^{-1} Mpc^{-3}}$ \\citep[comoving;][]{hoprich07,bong07} and adopting our estimate of \\mfp\\ at $z=4$, we find $\\mgq = 0.5 \\, \\mgigm$. This suggests a modest but non-negligible contribution from galaxies to the EUVB at this redshift.  The systematic uncertainties in $\\meq$ and \\gigm\\ are sufficiently large that one could recover \\gigm=\\gq, but \\gq\\ would overpredict \\gigm\\ at $z=3.5$ given the observed rise in \\mfp\\ and $\\meq$ with decreasing redshift.  The results also revise at least some previous estimates of the EUVB. \\cite{hm96} and their subsequent analyses (CUBA), for example, have adopted an approximately 3 times shorter mean free path at $z \\sim 4$ than our analysis reveals. This implies: (1) the normalization of their EUVB spectrum is several times too low; and (2) the EUVB spectrum is softer at energies of $\\approx 1$,\\,Ryd. The latter point may help to reconcile apparent contradictions in the metal-line analysis of the IGM \\citep{ads+08} without resorting to a large input from galaxies. Other analyses, however, have used estimates for \\mfp\\ that are in much better agreement with our results \\citep[e.g.][]{flz+09}. An understanding of the full implications of our new constraints on \\mfp\\ awaits new calculations of the EUVB.   We identified a previously unreported systematic bias in the SDSS quasar spectroscopic sample against sightlines at $z < 3.6$ that are `clear' of optically thick absorbers. This bias has implications for a range of IGM analysis at $z \\sim 3$ including: (i) the paucity of sightlines for studying \\ion{He}{2} reionization \\citep{wp09}; (ii) an overestimate of the incidence of LLS \\citep{pow09} and damped \\lya\\ systems \\citep{pw09}; and (iii) analysis of the \\lya\\ forest. We caution that all existing studies of the IGM at $z \\sim 3$ using the SDSS database should be reviewed in light of this systematic bias.  In future work, we analyze these same spectra to constrain the SEDs of $z > 3.5$ quasars, infer the \\nhi\\ frequency distribution and Doppler parameters of gas with $\\mnhi \\approx 10^{16} \\cm{-2}$, and isolate the opacity of the IGM far from the quasar's proximity region. The technique introduced here is easily extended to higher and lower redshifts by obtaining modest signal-to-noise (S/N), low-resolution spectroscopy of several hundred quasars. Future ground and space-based programs will precisely estimate \\kll\\ from $\\mzq \\approx 0.5 - 5$."
    },
    "0910/0910.1419_arXiv.txt": {
        "abstract": "The large scatters of luminosity relations of gamma-ray bursts (GRBs) have been one of the most important reasons that prevent the extensive applications of GRBs in cosmology. In this paper, we extend the two-dimensional (2D) luminosity relations with $\\tau_{\\mathrm{lag}}$, $V$, $E_{\\mathrm{peak}}$, and $\\tau_{\\mathrm{RT}}$ as the luminosity indicators to three dimensions (3D) using the same set of luminosity indicators to explore the possibility of decreasing the intrinsic scatters. We find that, for the 3D luminosity relations between the luminosity and an energy scale ($E_{\\mathrm{peak}}$) and a time scale ($\\tau_{\\mathrm{lag}}$ or $\\tau_{\\mathrm{RT}}$), their intrinsic scatters are considerably smaller than those of corresponding 2D luminosity relations. Enlightened by the result and the definition of the luminosity (energy released in units of time), we discussed possible reasons behind, which may give us helpful suggestions on seeking more precise luminosity relations for GRBs in the future. ",
        "introduction": "Gamma-ray bursts (GRBs) have recently attracted much attention in their cosmological applications as the most luminous astrophysical events observed today. Based on the correlations between the luminosity/energy and the measurable parameters of light curves and/or spectra, GRBs can be used as standard candles after calibration~\\citep[see, for example,][etc.]{Dai:2004tq, Ghirlanda:2004fs, Firmani:2005gs, Ghirlanda:2006ax, Schaefer:2006pa, Amati:2008hq, Basilakos:2008tp}. The advantage of the GRBs is their high redshifts due to their high luminosities. The $69$ GRBs compiled in~\\citet{Schaefer:2006pa} extend the redshift to $z>6$. Recently observed GRB 090423 has a redshift of $z=8.3$~\\citep{Tanvir:2009ex, Salvaterra:2009ey}. However, GRBs are not as ideal standard candles as type Ia supernovae (SNe Ia). Compared with SNe Ia, GRBs suffer from the circularity problem due to the lack of low-redshift samples and the scatters of known luminosity relations of GRBs are still very large. There are a few ways proposed in literatures to avoid the circularity problem. One can simply fit the calibration parameters and cosmological parameters simultaneously~\\citep{Li:2007re, Qi:2008zk}; and in~\\citet{Wang:2008vja}, GRB data are summarized by a set of model-independent distance measurements; calibrating GRBs using SNe Ia in their overlapping redshift range was also proposed~\\citep{Kodama:2008dq, Liang:2008kx} and adopted~\\citep{Wei:2008kq, Cardone:2009mr} in cosmological studies using GRBs.  Till now in most works of cosmological studies using GRBs, two-dimensional (2D) luminosity relations are used, which have least calibration parameters. For example $\\tau_{\\mathrm{lag}}$--$L$~\\citep{Norris:1999ni}, $V$--$L$~\\citep[][ there exist several definitions of $V$, mainly depending on the smoothing time intervals the reference curve is built upon, and on the normalization as well] {Fenimore:2000vs, Reichart:2000kq}, $E_{\\mathrm{peak}}$--$E_{\\gamma, \\mathrm{iso}}$~\\citep{Amati:2002ny}, $E_{\\mathrm{peak}}$--$E_{\\gamma}$~\\citep{Ghirlanda:2004me}, $E_{\\mathrm{peak}}$--$L$~\\citep{Schaefer:2002tf}, and $\\tau_{\\mathrm{RT}}$--$L$~\\citep{Schaefer:2006pa} relations. However, the intrinsic scatters of 2D luminosity relations are usually very large, which may imply hidden parameters considering the complication of GRBs. There are already works which explored the possibility of a three-dimensional (3D) correlation with negligible scatter. For example, \\citet{Firmani:2006gw} claimed that a temporal parameter of the prompt emission, the $T_{0.45}$, could reduce the scatter of the correlation of $L_{\\mathrm{iso}}$--$E_{\\mathrm{peak}}$ to a negligible value. But it was later found that the new proposed relation does not appear to be as tight as it seemed to be~\\citep{Rossi:2008mv, Collazzi:2008gu}. In this paper, we extend the luminosity relations used in~\\citet{Schaefer:2006pa} from 2D to 3D and explore the possibility of decreasing the scatters in the correlations. ",
        "conclusions": "\\begin{sidewaystable}[htbp] \\centering \\tiny \\begin{tabular}{|c|c|c|c|c|c|} \\hline (i, j) & 1 & 2 & 3 & 4 & 5 \\\\ \\hline 1 & \\begin{tabular}{@{} c @{}} $32$ \\\\ $(-3.994_{-0.077}^{+0.078}, \\, -0.79_{-0.11}^{+0.11})$ \\\\ $0.404_{-0.055}^{+0.067}$ \\end{tabular} & \\begin{tabular}{@{} c @{}} $22$ \\\\ $(-3.96_{-0.10}^{+0.10}, \\, -0.68_{-0.18}^{+0.19}, \\, 0.51_{-0.36}^{+0.38})$ \\\\ $0.414_{-0.067}^{+0.087}$ \\\\ $[-0.01_{-0.10}^{+0.09}, \\, 0.09_{-0.10}^{+0.09}]$ \\end{tabular} & \\begin{tabular}{@{} c @{}} $30$ \\\\ $(-4.013_{-0.060}^{+0.059}, \\, -0.618_{-0.091}^{+0.091}, \\, 0.80_{-0.16}^{+0.16})$ \\\\ $0.279_{-0.042}^{+0.052}$ \\\\ $[0.124_{-0.076}^{+0.079}, \\, 0.142_{-0.066}^{+0.064}]$ \\end{tabular} & \\begin{tabular}{@{} c @{}} $13$ \\\\ $(-5.575_{-0.083}^{+0.074}, \\, 0.08_{-0.14}^{+0.14}, \\, 1.44_{-0.18}^{+0.17})$ \\\\ $0.15_{-0.08}^{+0.11}$ \\\\ $[0.01_{-0.11}^{+0.10}]$ \\end{tabular} & \\begin{tabular}{@{} c @{}} $31$ \\\\ $(-3.952_{-0.081}^{+0.081}, \\, -0.60_{-0.14}^{+0.14}, \\, -0.31_{-0.17}^{+0.17})$ \\\\ $0.344_{-0.051}^{+0.063}$ \\\\ $[0.059_{-0.083}^{+0.084}, \\, 0.111_{-0.076}^{+0.072}]$ \\end{tabular} \\\\ \\hline 2 & - & \\begin{tabular}{@{} c @{}} $44$ \\\\ $(-3.712_{-0.081}^{+0.080}, \\, 1.02_{-0.22}^{+0.23})$ \\\\ $0.508_{-0.055}^{+0.066}$ \\end{tabular} & \\begin{tabular}{@{} c @{}} $42$ \\\\ $(-3.875_{-0.067}^{+0.067}, \\, 0.65_{-0.19}^{+0.19}, \\, 1.09_{-0.19}^{+0.19})$ \\\\ $0.366_{-0.043}^{+0.051}$ \\\\ $[0.141_{-0.075}^{+0.078}, \\, 0.055_{-0.065}^{+0.064}]$ \\end{tabular} & \\begin{tabular}{@{} c @{}} $22$ \\\\ $(-5.652_{-0.057}^{+0.052}, \\, 0.15_{-0.22}^{+0.23}, \\, 1.56_{-0.15}^{+0.15})$ \\\\ $0.178_{-0.051}^{+0.064}$ \\\\ $[-0.020_{-0.078}^{+0.075}]$ \\end{tabular} & \\begin{tabular}{@{} c @{}} $41$ \\\\ $(-3.702_{-0.071}^{+0.071}, \\, 0.46_{-0.23}^{+0.22}, \\, -0.65_{-0.16}^{+0.16})$ \\\\ $0.432_{-0.050}^{+0.060}$ \\\\ $[0.075_{-0.081}^{+0.082}, \\, 0.023_{-0.074}^{+0.071}]$ \\end{tabular} \\\\ \\hline 3 & - & - & \\begin{tabular}{@{} c @{}} $63$ \\\\ $(-3.999_{-0.058}^{+0.058}, \\, 1.37_{-0.12}^{+0.12})$ \\\\ $0.422_{-0.041}^{+0.048}$ \\end{tabular} & \\begin{tabular}{@{} c @{}} $27$ \\\\ $(-3.3_{-1.6}^{+1.6}, \\, 0.91_{-0.43}^{+0.48}, \\, 0.12_{-0.29}^{+0.27})$ \\\\ $0.50_{-0.14}^{+0.26}$ \\\\ $[-0.08_{-0.26}^{+0.14}, \\, -0.34_{-0.26}^{+0.15}]$ \\end{tabular} & \\begin{tabular}{@{} c @{}} $56$ \\\\ $(-3.852_{-0.048}^{+0.048}, \\, 0.90_{-0.11}^{+0.11}, \\, -0.624_{-0.090}^{+0.089})$ \\\\ $0.293_{-0.032}^{+0.037}$ \\\\ $[0.128_{-0.055}^{+0.057}, \\, 0.162_{-0.057}^{+0.060}]$ \\end{tabular} \\\\ \\hline 4 & - & - & - & \\begin{tabular}{@{} c @{}} $27$ \\\\ $(-5.626_{-0.047}^{+0.044}, \\, 1.51_{-0.11}^{+0.11})$ \\\\ $0.159_{-0.046}^{+0.054}$ \\end{tabular} & \\begin{tabular}{@{} c @{}} $23$ \\\\ $(-5.661_{-0.059}^{+0.055}, \\, 1.56_{-0.14}^{+0.15}, \\, 0.00_{-0.12}^{+0.11})$ \\\\ $0.176_{-0.052}^{+0.064}$ \\\\ $[-0.017_{-0.078}^{+0.075}]$ \\end{tabular} \\\\ \\hline 5 & - & - & - & - & \\begin{tabular}{@{} c @{}} $61$ \\\\ $(-3.766_{-0.065}^{+0.065}, \\, -0.88_{-0.11}^{+0.11})$ \\\\ $0.456_{-0.044}^{+0.051}$ \\end{tabular} \\\\ \\hline \\end{tabular} \\caption { Fit of 2D and 3D luminosity relations. In every grid, the first row is the number of GRBs of set $(i, \\, j)$, the vector below enclosed by parentheses is the vector of $c$ for the luminosity relation $(i, \\, j)$, and what follows next is the intrinsic scatter. For 3D luminosity relations, their reduction in the intrinsic scatters compared to corresponding 2D luminosity relations are presented in the brackets: for 3D luminosity relations in Eq.~(\\ref{eq:GRB-3D_1}) and Eq.~(\\ref{eq:GRB-3D_2}), the reduction corresponds to the 2D luminosity relation $(i, \\, i)$ and $(j, \\, j)$ in turn and for those in Eq.~(\\ref{eq:GRB-3D_3}), corresponds to $(4, \\, 4)$. The statistics in the table are for the median values and the errors of $1 \\sigma$ ($68.3\\%$) confidence level. } \\label{tab:fit} \\end{sidewaystable}  \\begin{sidewaystable}[htbp] \\centering \\begin{tabular}{|c|c|c|c|c|c|c|} \\hline & $x^{(1)}$ & $x^{(2)}$ & $x^{(3)}$ & $x^{(5)}$ & $y^{(3)}$ & $y^{(4)}$ \\\\ \\hline $(1, \\, 1)$ & - & $0.31_{-0.14}^{+0.12}$ & $0.712_{-0.077}^{+0.053}$ & $-0.18_{-0.14}^{+0.15}$ & - & - \\\\ \\hline $(2, \\, 2)$ & $-0.647_{-0.069}^{+0.093}$ & - & $0.689_{-0.060}^{+0.042}$ & $-0.474_{-0.081}^{+0.097}$ & - & - \\\\ \\hline $(3, \\, 3)$ & $-0.732_{-0.040}^{+0.046}$ & $0.472_{-0.044}^{+0.038}$ & - & $-0.634_{-0.044}^{+0.050}$ & - & $-0.21_{-0.09}^{+0.10}$ \\\\ \\hline $(4, \\, 4)$ & $0.223_{-0.012}^{+0.008}$ & $0.380_{-0.020}^{+0.014}$ & - & $-0.007_{-0.037}^{+0.037}$ & $-0.09_{-0.11}^{+0.11}$ & - \\\\ \\hline $(5, \\, 5)$ & $-0.36_{-0.10}^{+0.10}$ & $0.185_{-0.089}^{+0.087}$ & $0.740_{-0.045}^{+0.034}$ & - & - & - \\\\ \\hline \\end{tabular} \\caption { Correlation coefficients between the residuals of the fit of 2D correlations and possible hidden parameters. The statistics in the table are for the median values and the errors of $1 \\sigma$ ($68.3\\%$) confidence level. } \\label{tab:correlation} \\end{sidewaystable}  We summarize our results in Tables~\\ref{tab:fit} and~\\ref{tab:correlation}. In Table~\\ref{tab:fit}, by comparing the intrinsic scatters of the 3D luminosity relations with those of the corresponding 2D luminosity relations, we find that only for the cases of $(1, \\, 3)$ and $(3, \\, 5)$ the intrinsic scatters of the 3D luminosity relations are considerably smaller than those of their corresponding 2D luminosity relations, and for all other cases, the intrinsic scatters of the 3D luminosity relations are either very close to the smaller one of the intrinsic scatters of the corresponding 2D luminosity relations (for correlations in Eq.~(\\ref{eq:GRB-3D_3}), close to the intrinsic scatter of luminosity relation $(4, \\, 4)$) or even greater than those of the corresponding 2D luminosity relations (for the case $(3, \\, 4)$). The correlations presented in Table~\\ref{tab:correlation} also give the same implication. Only the correlation coefficients between the residual of fit of $(1, \\, 1)$ and $x^{(3)}$, $(3, \\, 3)$ and $x^{(1)}$, $(5, \\, 5)$ and $x^{(3)}$ are approximately greater than $0.7$. We presented the plots of the luminosity relation $(1, \\, 3)$ and $(3, \\, 5)$ in Figure~\\ref{fig:xy_plot}. \\begin{figure}[htbp] \\centering \\includegraphics[width = 0.45 \\textwidth]{xy13.eps} \\includegraphics[width = 0.45 \\textwidth]{xy35.eps} \\caption { Plots of the 3D luminosity relation $(1, \\, 3)$ and $(3, \\, 5)$. } \\label{fig:xy_plot} \\end{figure}  It is interesting to note that the 3D luminosity relations that have considerable improvement in the intrinsic scatter are the luminosity relations in Eq.~(\\ref{eq:GRB-3D_1}), with one of the luminosity indicators being an energy scale ($E_{\\mathrm{peak}}$) and the other a time scale ($\\tau_{\\mathrm{lag}}$ or $\\tau_{\\mathrm{RT}}$). Enlightened by this, one may naturally guess that probably the luminosities of GRBs are mainly determined by a characteristic energy scale and a characteristic time scale and $E_{\\mathrm{peak}}$ is correlated to the characteristic energy scale, $\\tau_{\\mathrm{lag}}$ and $\\tau_{\\mathrm{RT}}$ to the characteristic time scale, so that 3D luminosity relations with $E_{\\mathrm{peak}}$ and $\\tau_{\\mathrm{lag}}$ or $\\tau_{\\mathrm{RT}}$ as luminosity indicators could significantly reduce the intrinsic scatters compared to corresponding 2D luminosity relations. This is easy to understand, since the definition of the luminosity is the amount of the energy released in units of time. If one knows the released energy and the time duration of a GRB, the (averaged) luminosity can be calculated immediately. Obviously, the 2D correlation between the luminosity and the energy or the time duration would not be complete unless the time duration or the energy is constant for all samples or there is a strong correlation between the energy and the time duration. This, if proven true, also explains why the 2D luminosity relations $(3, \\, 3)$ and $(4, \\, 4)$ have little correlation with each other despite their similarity at first sight (see discussion in Section 4.7 of~\\citet{Schaefer:2006pa}; the correlation coefficient between the residual of fit of $(3, \\, 3)$ and $y^{(4)}$ and the correlation coefficient between the residual of fit of $(4, \\, 4)$ and $y^{(3)}$ also show their weak correlation). This is because of the difference between $E_{\\gamma}$ and $L$, where a new independent variable---the characteristic time scale---enters. In fact, we could check the guess partially with current data by examining the correlation between $x^{(i)}$s. In Table~\\ref{tab:cor_x}, one can find that $x^{(1)}$, $x^{(2)}$, and $x^{(5)}$ are strongly correlated with each other, while all of them have only weak correlation with $x^{(3)}$ (see also~\\citet{Schaefer:2001ap} and Section 4.5 of~\\citet{Schaefer:2006pa} for discussions on the correlations between the luminosity indicators). This is consistent with the guess that the luminosity indicators $\\tau_{\\mathrm{lag}}$ and $\\tau_{\\mathrm{RT}}$ are correlated with a characteristic time scale and $E_{\\mathrm{peak}}$ is correlated with a characteristic energy scale which is independent of the characteristic time scale. In addition, the correlations presented in Table~\\ref{tab:cor_x} seems to imply that $V$ is also correlated with the characteristic time scale. Accordingly, there is indeed some reduction in the intrinsic scatter for the 3D luminosity relation $(2, \\, 3)$ compared with corresponding 2D luminosity relations, but such reduction is relatively smaller than that of $(1, \\, 3)$ and $(3, \\, 5)$. Maybe the correlation of $V$ with the characteristic time scale is not so strong as that of $\\tau_{\\mathrm{lag}}$ or $\\tau_{\\mathrm{RT}}$. \\begin{table}[htbp] \\centering \\begin{tabular}{|c|c|c|c|c|} \\hline &          $x^{(1)}$ & $x^{(2)}$ & $x^{(3)}$ & $x^{(5)}$ \\\\ \\hline $x^{(1)}$ & 1.0      & -0.73    & -0.35    & 0.72 \\\\ \\hline $x^{(2)}$ & -0.73    & 1.0      & 0.41     & -0.63 \\\\ \\hline $x^{(3)}$ & -0.35    & 0.41     & 1.0      & -0.30 \\\\ \\hline $x^{(5)}$ & 0.72     & -0.63    & -0.30    & 1.0 \\\\ \\hline \\end{tabular} \\caption { Correlation coefficients between $x^{(1)}$, $x^{(2)}$, $x^{(3)}$, and $x^{(5)}$. } \\label{tab:cor_x} \\end{table}  If our guess about the GRB luminosity relations is correct, it would be very enlightening. It suggests us to include an energy scale and a corresponding time scale for seeking more precise luminosity relations for GRBs in the future. However, it should be emphasized that, even if it is true, only appropriate energy and time scales might significantly reduce the intrinsic scatter bearing in mind the situation of some 3D luminosity relations~\\citep[see, for example,][]{Firmani:2006gw, Rossi:2008mv, Collazzi:2008gu}."
    },
    "0910/0910.1133_arXiv.txt": {
        "abstract": "The gravitational lens system CLASS B2108+213 has two radio-loud lensed images separated by 4.56 arcsec. The relatively large image separation implies that the lensing is caused by a group of galaxies. In this paper, new optical imaging and spectroscopic data for the lensing galaxies of B2108+213 and the surrounding field galaxies are presented. These data are used to investigate the mass and composition of the lensing structure. The redshift and stellar velocity dispersion of the main lensing galaxy (G1) are found to be $z=$~0.3648\\,$\\pm$\\,0.0002 and $\\sigma_v =$~325\\,$\\pm$\\,25~km\\,s$^{-1}$, respectively. The optical spectrum of the lensed quasar shows no obvious emission or absorption features and is consistent with a BL Lac type radio source. However, the tentative detection of the G-band and Mg-b absorption lines, and a break in the spectrum of the host galaxy of the lensed quasar gives a likely source redshift of $z=$~0.67. Spectroscopy of the field around B2108+213 finds 51 galaxies at a similar redshift to G1, thus confirming that there is a much larger structure at $z\\sim$~0.365 associated with this system. The width of the group velocity distribution is 694\\,$\\pm$\\,93~km\\,s$^{-1}$, but is non-Gaussian, implying that the structure is not yet viralized. The main lensing galaxy is also the brightest group member and has a surface brightness profile consistent with a typical cD galaxy. A lensing and dynamics analysis of the mass distribution, which also includes the newly found group members, finds that the logarithmic slope of the mass density profile is on average isothermal inside the Einstein radius, but steeper {\\it at} the location of the Einstein radius. This apparent change in slope can be accounted for if an external convergence gradient, representing the underlying parent halo of the galaxy group, is included in the mass model. ",
        "introduction": "In the standard model of galaxy formation, structure arises hierarchically through the mergers and interactions of smaller mass systems, leading to increasingly more massive structures. Cold dark matter (CDM) simulations of these merging events predict universal radial matter density profiles with inner slopes of $\\gamma \\sim$~1--1.5 [where $\\rho_{\\rm CDM}(r) \\propto r^{-\\gamma}$], which are independent of the halo mass [e.g. \\citealt*{navarro96} (NFW); \\citealt{moore98}; \\citealt*{diemand07,navarro09}]. Furthermore, the compact cores of merging systems are expected to survive the cannibalism process, forming CDM substructure within the new parent halo (e.g. \\citealt{moore99,gao04,diemand07}).  Strong gravitational lensing, where multiple images of the same background object are formed by an intervening potential, has provided a powerful observational test of these theoretical predictions for mass structures out to $z \\sim$~1, particularly when combined with stellar kinematic data (e.g. \\citealt{treu04,sand04}; \\citealt{koopmans06}).  In the case of galaxy-scale lensing ($M_{\\rm lens}=$\\,10$^{11}$--10$^{12}~$M$_{\\odot}$), detailed mass modeling of the positions and flux densities of the multiply imaged source has shown the total (i.e. baryons plus dark matter; $\\rho_{\\rm tot}$) mass density profiles of lens galaxies to be close to isothermal, that is, $\\gamma'=d \\log \\rho_{\\rm tot}/d\\log r \\sim$~2 (\\citealt{cohn01}; \\citealt{koopmans06}; \\citealt*{rusin03a,treu04,wucknitz04}; \\citealt{koopmans09}) with very little scatter ($\\sim$5 per cent on $\\gamma'$).  The steeper slope than that predicted by dark matter only simulations is generally attributed to baryon cooling and collapse, which steepens the inner density profile for lens galaxies \\citep{blumenthal84,gnedin04,keeton01b,kochanek01a,jiang07}. The origin of the small scatter in the total mass density profile slope is not completely clear (see discussion in \\citealt{koopmans06}).  Decomposing the baryonic and dark matter components is particularly hard for galaxy-scale lenses, given the dominance of baryonic mass at small radii. By combining lensing with stellar kinematic tracers, and assuming a spatially uniform mass-to-light ratio for the stellar distribution, an upper limit of $\\gamma \\sim$~1.3 for the dark matter component of galaxy-scale haloes \\citep{treu04} can be obtained, which is certainly within the allowable range predicted from simulations. However, studies of the inner slopes of dark matter haloes using non-lensing techniques, for example, from rotation curve measurements of nearby dwarf and spiral galaxies, find evidence for a flattened core (0 $< \\gamma <$ 1) out as far as the optical radius (e.g. \\citealt{salucci00,swaters03,gentile04,donato09}), which are typically inconsistent with the CDM models. Although, it should be noted that due to systematic effects, such as non-circular motions or poorly defined disk inclinations, the CDM models cannot be discounted \\citep{swaters03}.  For galaxy clusters ($M_{\\rm lens}>$\\,10$^{15}$~M$_{\\odot}$), the situation is significantly different. On the one hand, the stellar component is less important than for galaxies on the scales that are relevant to measure the inner slope of the dark matter, and several mass probes are available to cover the broad range of radii needed for this measurement, including strong lensing, stellar kinematics, X-rays and weak lensing. On the other hand, many clusters are not dynamically relaxed systems, often with significant triaxiality and substructure. Therefore the interpretation of the results can be more complicated than in the case of galaxies and can be affected by potential systematics, such as those arising from projection effects and departures from hydrostatic equilibrium.  A relatively large scatter of inner density profile slopes has been reported for galaxy clusters (\\citealt{smith01}; \\citealt*{sand02}; \\citealt*{ettori02}; \\citealt*{lewis03}; \\citealt{sand04,sand05}; \\citealt{gavazzi05}; \\citealt*{saha06}; \\citealt{sharon05}; \\citealt{bradac08}) -- ranging from flat inner cores to slopes consistent with CDM predictions. However, it should be noted that some of these authors refer to the slope of the dark matter component ($\\gamma$) or to the slope of the total mass density profile ($\\gamma'$) as if they were the same quantity. This is not justified. In fact, although stars are only a small fraction of the total mass, they are extremely spatially concentrated and therefore they affect the {\\it slope} in a significant way (\\citealt{sand02}). Overall, although more work remains to be done at cluster scales, it seems clear that in at least some cases the inner slope of the dark matter halo is flatter than predicted by CDM simulations and the NFW profile is not a good fit to the data (\\citealt{sand08}).  Galaxy groups ($M_{\\rm lens}=$\\,$10^{13}$--$10^{14}~$M$_{\\odot}$), containing a few to a few tens of galaxies, are the most common environment for galaxies to reside at low redshifts (e.g. \\citealt{eke04}). Although the total mass of a galaxy group can be estimated from the luminosity distribution or the stellar mass of the individual group members (e.g. \\citealt{yang07}) as well as from the group velocity dispersion \\citep{zabludoff98}, the distribution of the matter within groups is not well known.  Galaxy groups have been relatively unexplored with strong gravitational lensing due to the paucity of known direct image splitting by group sized haloes (see \\citealt{limousin09} and \\citealt{kubo09} for recent developments) and weak lensing measurements can require the stacking of possibly hundreds to thousands of groups (e.g. \\citealt{hoekstra01,parker05,mandelbaum06}), and even then, only probe the outer part of the dark matter halo. There are several examples where strong gravitational lensing galaxies have been found to be associated with a group of galaxies and/or there are groups of galaxies found along the line-of-sight to the lensed object \\citep{tonry98,tonry99,tonry00,fassnacht02,faure04,morgan05,fassnacht06,momcheva06,williams06,auger07}. In these cases, the lensed image separations are still in the galaxy-scale regime of $\\sim$\\,0.5--1.5 arcsec because the lensing potential is dominated by a single galaxy, with the contribution of the group to the overall convergence being only $\\sim$5 per cent \\citep{momcheva06,auger07,auger08}. However, it should be noted that environmental contributions at this level can still have a large impact on the results of lens modelling \\citep{keeton04}. For lens systems with image separations larger than 3~arcsec the contribution of the environment is thought to be much larger than in the galaxy-scale regime (with an external convergence $\\ga$~10 per cent), possibly enhancing the lensing probabilities \\citep*{oguri05}.  How the mass is distributed within a group-sized halo, and what effect baryons have on the overall density profile is not entirely clear. However, as the image separation increases there is expected to be a transition from galaxy scale haloes, which are dominated by stars in the centre and have a relatively steep total mass density profile (i.e. $\\gamma' \\sim$~2), to group-scale haloes, characterized by a flatter inner density profile of $\\gamma' \\sim$~1--1.5 \\citep{oguri06}. Another question to be addressed is whether the dense environment affects the dark matter sub-haloes of the satellite galaxies within the group. It is expected that due to tidal interactions, the dark matter of the sub-haloes will be stripped, resulting in a change in the sub-halo mass-to-light ratio (e.g. \\citealt{gao04}) and a truncation which would lead to a more concentrated sub-halo mass distribution (e.g. \\citealt{bullock01}).  In this paper, we present for the first time an analysis of the total (baryonic and dark) inner density profile of a moderate redshift galaxy group using strong gravitational lensing and stellar kinematics. We show that the gravitational lens system B2108+213 is part of a massive group of galaxies and that the multiple imaging is found around the large central galaxy, as opposed to the satellite members of the group. Thus, we are probing the inner part of the group-scale halo. In Section \\ref{classb2108+213}, we present an introduction to the lens system and summarize the previous observations. In Section \\ref{photometry}, new ground based imaging with the W. M. Keck Telescope is combined with archival {\\it HST} data to select potential members of the galaxy group associated with B2108+213. Our spectroscopy of group candidates is presented in Section \\ref{spectroscopy}. From these observations we found 51 galaxies at the same redshift as the main lensing galaxy. The properties of the group and the constituent group members are discussed in Section \\ref{redshifts}. An analysis of the stellar velocity dispersion and surface brightness profile of the main lensing galaxy is presented in Section \\ref{kinematics}. In Section \\ref{modelling} we present a new mass model for B2108+213 which incorporates the group galaxies and the stellar kinematics. We discuss our results in Section \\ref{disc} and present our conclusions in Section \\ref{conc}.  For all calculations we adopt an $\\Omega_{M} =$~0.3, $\\Omega_{\\Lambda}=$~0.7 spatially-flat Universe, with a Hubble constant of $H_{0} =$~70~km\\,s$^{-1}$~Mpc$^{-1}$. At redshift 0.365, 1 arcsec corresponds to an angular size of 5.075~kpc. ",
        "conclusions": "\\label{conc}  We have presented new optical imaging and spectroscopy of the gravitational lens system B2108+213. Our observations have shown that the main lensing galaxy is the central galaxy of a massive group or poor cluster. Using both gravitational lensing and stellar kinematical analyses, we have carried out extensive modelling of the system and have come to the following conclusions.  i) By targeting a gravitational lens system with a wide image separation, we have successfully identified a large group of galaxies at a moderate redshift. Recent studies of the environments of gravitational lenses have found many are associated with groups of galaxies. However, these systems tend to have lensed image separations of order $\\sim$\\,1--1.5~arcsec, and are therefore, typically associated with poor groups, giving only limited information on the mass distribution of dark matter haloes in the mass regime between galaxies and massive clusters. Further investigations of group-scale dark matter haloes should target gravitational lens systems with image separations similar to or larger than $\\sim$~4~arcsec.  ii) Our modelling of the inner part of the group mass distribution has found that the combined dark and baryonic matter density profile is much steeper (i.e. $\\gamma' \\ga$~2) when compared to what is predicted from dark-matter only simulations (i.e. $\\gamma \\sim$~1--1.5). This suggests that in the case of the group associated with B2108+213, the baryonic component within the inner 20~kpc of the halo centre dominates the mass distribution. B2108+213 may be a special case in this respect since there is a massive cD galaxy in the centre of the group. Therefore, it would be useful to find and study group systems without a massive central galaxy in the future.  iii) Our analysis of the gravitational lensing and stellar dynamics data finds that the inner matter density profile is on average isothermal within the Einstein radius, but is steeper (i.e. $\\gamma' \\sim$~2.5) at the position of the Einstein radius. A plausible explanation for this contradiction is the presence of an external convergence gradient across the lens system. The effect of this gradient is to change the flux-ratio of the lensed images and to bring the slope of the inner matter density profile determined from lensing to be consistent with what is predicted from the stellar dynamics of the lensing galaxy. The presence of a convergence gradient is certainly plausible considering there is a galaxy group associated with G1 and G2. Therefore, we find direct evidence for the environment of the lens system having an effect on the observed properties of the lensed images. Note that by not including a convergence gradient, the inner matter density profile is artificially steepened and this should be considered when modelling other group-scale systems in the future.  Our conclusions are based on the results from investigating only a single group system with gravitational lensing. Clearly, studies of many more group systems will be needed to test our findings for the B2108+213 group. Recently, there have been 13 group-scale gravitational lens systems identified from the Strong Lensing Legacy Survey (S2LS; \\citealt{limousin09}). These groups span a large range of group masses and, apparently from maps of the luminosity distributions, a large scatter of group halo morphologies. It will be interesting to see whether this sample of strong lensing groups and poor clusters can be modelled without including an external convergence gradient."
    },
    "0910/0910.2158.txt": {
        "abstract": "{The carbon-enhanced metal-poor (CEMP) stars constitute approximately one fifth of the metal-poor ($[\\mathrm{Fe}/\\mathrm{H}]\\lesssim-2$) population but their origin is not well understood. The most widely accepted formation scenario, at least for the majority of CEMP stars which are also enriched in $s$-process elements, invokes mass-transfer of carbon-rich material from a thermally-pulsing asymptotic giant branch (TPAGB) primary star to a less massive main-sequence companion which is seen today. Recent studies explore the possibility that an initial mass function biased toward intermediate-mass stars is required to reproduce the observed CEMP fraction in stars with metallicity $[\\mathrm{Fe}/\\mathrm{H}]<-2.5$. These models also implicitly predict a large number of nitrogen-enhanced metal-poor (NEMP) stars which is not seen. In this paper we investigate whether the observed CEMP and NEMP to extremely metal-poor (EMP) ratios can be explained \\emph{without} invoking a change in the initial mass function. We construct binary-star populations in an attempt to reproduce the observed number and chemical abundance patterns of CEMP stars at a metallicity $[\\mathrm{Fe}/\\mathrm{H}]\\sim-2.3$. Our binary-population models include synthetic nucleosynthesis in TPAGB stars and account for mass transfer and other forms of binary interaction. This approach allows us to explore uncertainties in the CEMP-star formation scenario by parameterization of uncertain input physics. In particular, we consider the uncertainty in the physics of third dredge up in the TPAGB primary, binary mass transfer and mixing in the secondary star. We confirm earlier findings that with current detailed TPAGB models, in which third dredge up is limited to stars more massive than about $1.25\\mathrm{\\, M_{\\odot}}$, the large observed CEMP fraction cannot be accounted for. We find that efficient third dredge up in low-mass (less than $1.25\\mathrm{\\, M_{\\odot}}$), low-metallicity stars may offer at least a partial explanation to the large observed CEMP fraction while remaining consistent with the small observed NEMP fraction. } ",
        "introduction": "\\label{sec:Introduction}One of the most interesting problems in modern stellar astronomy is to explain the existence of a population of carbon-enhanced metal-poor (CEMP% \\footnote{$[\\mathrm{C}/\\mathrm{Fe}]\\geq1$, $\\mathrm{\\left[\\mathrm{Fe}/\\mathrm{H}\\right]}\\lesssim-2$, where the logarithmic abundance ratio $[X/Y]=\\log_{10}(X/Y)-\\log_{10}(X_{\\odot}/Y_{\\odot})$.% }) stars in the Galactic halo. The HK \\citep{1992Beers++} and Hamburg/ESO \\citep{2001A&A...375..366C} surveys find a large number of CEMP stars among the metal-poor (EMP, $[\\mathrm{Fe}/\\mathrm{H}]\\lesssim-2$) population, at a fraction around 20\\% (e.g. $9\\pm2\\%$ \\citealp{2006ApJ...652.1585F}; $21\\pm2\\%$ \\citealp{2006ApJ...652L..37L}; up to $30\\%$ from the SAGA database of \\citealp{2008PASJ...60.1159S} -- see below for details).  The CEMP stars are subdivided into four groups depending on the presence or absence of the heavy elements barium and europium (see e.g. \\citealp{2005ARA&A..43..531B}). The most populous group consists of the $s$-process rich CEMP stars, the so-called CEMP-$s$ stars (e.g. \\citealp{2007ApJ...655..492A}) which display barium enhancements of $[\\mathrm{Ba}/\\mathrm{Fe}]>+0.5$. These account for about 80 per cent of all CEMP stars. There are also CEMP stars with $r$-process enhancements (the CEMP-$r$ class) and some with both $r$- and $s$-process enhancements (CEMP-$r$+$s$, e.g. \\citealp{2006A&A...451..651J}). Finally, there is a class of CEMP stars which show no enhancement of neutron-capture elements. These are called the CEMP-no stars \\citep{2002Aoki++_no_ncapt}. A detailed review of the various CEMP subgroups can be found in \\citet{2009arXiv0901.4737M}.  A quantitative understanding of the origin of CEMP stars touches on many branches of stellar astronomy. The most likely formation mechanism for the $s$-process rich CEMP stars involves mass transfer in binary systems. Carbon-rich material from the TPAGB primary star pollutes the lower-mass main sequence secondary such that it becomes enriched in carbon and $s$-process elements. We observe only the secondary today; the primary is an unseen white dwarf. Surveys of radial velocity shifts find that the binary fraction of CEMP-$s$ stars is consistent with them all being binaries \\citep{2004MmSAI..75..772T,2005ApJ...625..825L}. This binary mass transfer scenario is the same as that which is invoked to explain the Ba and CH stars \\citep{1983ARA&A..21..271I,1984ApJ...280L..31M,1990mcclure_woodsworth,1997mcclure}. However, only about $1\\%$ of population I/II stars are Ba/CH stars, respectively \\citep{1989A&A...219L..15T,1991ApJS...77..515L}. The CH stars are also carbon-rich but not as metal-poor as the CEMP stars, with $[\\mathrm{Fe}/\\mathrm{H}]\\sim-1$. The carbon-rich fraction of 1 per cent at higher metallicity is in stark contrast to the observed CEMP fraction of around 20 per cent.  The mass-transfer scenario involves many processes that are not well understood. There are uncertainties associated with stellar evolution, particularly with respect to nucleosynthesis in TPAGB stars. Carbon and $s$-process enhancements are thought to occur via third dredge up, but other processes may also play an important role in low-metallicity nucleosynthesis. These include hot-bottom burning (e.g. \\citealp{1975ApJ...196..525I,1993ApJ...416..762B,2004Herwig}), dual core flashes (also known as helium flash driven deep mixing) and dual shell flashes (helium flash driven deep mixing during a thermal pulse, \\citealt*{1990ApJ...349..580F,2002Schlattl++,2007ApJ...667..489C,2008A&A...490..769C}), {}``extra mixing'' on the first giant branch and perhaps on the TPAGB (e.g. \\citealt*{2000A&A...356..181W,2003ApJ...582.1036N,2008ApJ...677..581E}) and convective overshooting \\citep{2000A&A...360..952H}.  The CEMP-formation scenario also requires knowledge of the physics of stellar interaction in binary systems. The primary star must transfer material to the secondary star which in turn might dilute and burn it. A wind mass-transfer scenario in wide binaries, e.g. by a mechanism similar to that of \\citet{1944MNRAS.104..273B}, likely plays a role in CEMP formation. Closer binaries which undergo Roche-lobe overflow (RLOF) from a TPAGB star on to a less massive main-sequence star are expected to enter a common-envelope phase. Such stars would undergo few thermal pulses with little accretion on to the secondary (although see \\citealp{2008ApJ...672L..41R} for details of accretion in a common envelope and associated uncertainties).  The fate of the accreted material is also uncertain. The molecular weight of accreted material is certainly greater than that of the secondary star, as carbon-enhanced TPAGB stars should also be helium rich. Accreted material should thus sink by the thermohaline instability \\citep{2007A&A...464L..57S} but this may be inhibited by gravitational settling \\citep{2008MNRAS.389.1828S,2008ApJ...677..556T}. Furthermore, radiative levitation of some chemical species may be important \\citep{2002ApJ...580.1100R,2002ApJ...568..979R}. When the secondary ascends the first giant branch, its convection zone mixes any accreted material which may remain in the surface layers with material from deep inside the star. The surface abundance distribution depends on whether material has mixed deep into the star or not, because if it has it may have undergone some nuclear burning, the ashes of which are mixed to the surface.  Two studies have considered population models in an attempt to reproduce the observed CEMP fraction of about $20\\%$. The models of \\citet{2005ApJ...625..833L} and \\citet{2007ApJ...658..367K} both concluded that in order to make enough CEMP stars the initial mass function (IMF) at low metallicity must be significantly different to that observed in the solar neighbourhood. In particular, they enhanced the number of intermediate-mass stars relative to low-mass stars -- this has the effect of increasing the number of $\\sim2\\mathrm{\\, M_{\\odot}}$ stars which are responsible for the production of most of the carbon.  The \\citet{2007ApJ...658..367K} model differentiates between two metallicity regimes. In their model, stars with masses greater than $1.5\\mathrm{\\, M_{\\odot}}$ undergo third dredge up irrespective of metallicity. For stars with $[\\mathrm{Fe}/\\mathrm{H}]\\leq-2.5$ and mass less than $1.5\\mathrm{\\, M_{\\odot}}$ they invoke proton ingestion at the helium flash (the dual core flash) or at the first thermal pulse (the dual shell flash) as the source of carbon. This implies that the CEMP fraction should be smaller for $[\\mathrm{Fe}/\\mathrm{H}]>-2.5$, but the Stellar Abundances for Galactic Archeology (SAGA) database \\citep{2008PASJ...60.1159S} shows that the CEMP fraction is approximately constant as a function of metallicity up to $[\\mathrm{Fe}/\\mathrm{H}]\\approx-2$.  A related problem is that of the nitrogen-enhanced metal-poor (NEMP) stars which have $[\\mathrm{N}/\\mathrm{Fe}]>0.5$ and $[\\mathrm{C}/\\mathrm{N}]<-0.5$ as defined by \\citet{2007ApJ...658.1203J}. Such stars are expected to result from mass transfer in binaries with TPAGB primaries more massive than about $3\\, M_{\\odot}$ in which hot bottom burning has converted most of the dredged-up carbon into nitrogen. The observed NEMP to EMP ratio is small, less than one in twenty-one according to \\citet{2007ApJ...658.1203J} or less than $7\\%$ in the metallicity range $[\\mathrm{Fe}/\\mathrm{H]=-2.3\\pm0.5}$ according to the SAGA database (see Table~\\ref{tab:Observations-table}). The \\citet{2007ApJ...658..367K} models, with an enhanced number of intermediate-mass relative to low-mass stars, should make many more NEMP stars than are observed.  The aim of this paper is to investigate which physical scenarios are able to reproduce the CEMP and NEMP to EMP ratios, at metallicity $[\\mathrm{Fe}/\\mathrm{H}]\\sim-2.3$, \\emph{without} altering the initial mass function. We combine a synthetic nucleosynthesis model with a binary population synthesis code to simulate populations of low-metallicity binaries (see Section~\\ref{sec:Model}). The power of the population synthesis approach is that it can efficiently explore the available parameter space. Much of the input physics is uncertain (as we have described above) but population synthesis allows us to explore the consequences of these uncertainties by varying the model free parameters within reasonable bounds. We try to reproduce the observed CEMP and NEMP to EMP ratios, surface chemistry distributions, binary period distributions and chemical abundance correlations. It is thus a powerful tool to apply to this problem.  In order to compare our models to observations we choose a subset of the SAGA database which corresponds to giants and turn-off stars as described in Section~\\ref{sec:Observational-database}. The results of our simulations and comparison with the sample of observations are given in Section~\\ref{sec:Results}. The implications of our results and outstanding problems are discussed in Section~\\ref{sec:Discussion} while Section~\\ref{sec:Conclusions} concludes. ",
        "conclusions": "In an attempt to reproduce the observed CEMP to EMP number ratio we have simulated populations of low-metallicity ($[\\mathrm{Fe}/\\mathrm{H}]=-2.3$) binary stars with a variety of input physics. Our model sets with efficient third dredge up in low-mass (down to $0.8\\mathrm{\\, M_{\\odot}}$) stars have CEMP to EMP ratios of up to $15\\%$, comparable with the observed $\\sim20\\%$. They also have low NEMP to EMP number ratios, in agreement with the observations. Other parameters in our simulations, such as the efficiency of wind accretion, the common envelope parameter etc., have only relatively minor effects on our results.  \\label{sec:Conclusions}"
    },
    "0910/0910.3246_arXiv.txt": {
        "abstract": "Recent investigations seem to favor a cosmological dynamics according to which the accelerated expansion of the Universe may have already peaked and is now slowing down again \\cite{sastaro}. As a consequence, the cosmic acceleration may be a transient phenomenon. We investigate a toy model that reproduces such a  background behavior  as the result of a time-dependent coupling in the dark sector which implies a cancelation of the ``bare\" cosmological constant. With the help of a statistical analysis of Supernova Type Ia (SNIa) data we demonstrate that for a certain parameter combination a transient accelerating phase emerges as a pure interaction effect. ",
        "introduction": "By now there is large direct and indirect evidence that the  Universe entered a period of accelerated expansion. Direct evidence is provided by the luminosity-distance data of supernovae of type Ia \\cite{SNIa} (see, however, \\cite{sarkar}), indirect evidence comes from the anisotropy spectrum of the cosmic microwave background radiation (CMBR) \\cite{cmb}, from large-scale-structure data \\cite{lss}, from the integrated Sachs--Wolfe effect \\cite{isw}, from baryonic acoustic oscillations \\cite{eisenstein} and from gravitational lensing \\cite{weakl}. According to the currently prevailing interpretation, based on Einstein's GR,  our Universe is dynamically dominated by two so far unknown components, dark matter (DM) and dark energy (DE). The latter contributes roughly 75\\% to the total energy budget, the former about 20\\%. Only about 5\\% are in the form of conventional, baryonic matter. The preferred model is the $\\Lambda$CDM model which also plays the role of a reference model for alternative approaches to the DE problem. While the $\\Lambda$CDM model can describe most of the observations, there still remain puzzles \\cite{peri}. The present situation in the field is characterized, e.g., in \\cite{rev,pad,dumarev} and references therein.  Both DM and DE manifest themselves observationally only through their gravitational action. In most studies both these components are regarded as independent. However, there exists a line of investigation that explores the consequences of a coupling within the dark sector. Corresponding models show a richer dynamics than models with separately conserved energy-momentum tensors. It has been argued that ignoring a potentially existing interaction between both components may give rise to misinterpretations of observational data regarding the equation-of-state (EoS) parameter \\cite{Das}. A coupling between DM and DE may also be relevant with respect to the coincidence problem (see, e.g. \\cite{coinc} and references therein). Models with an interaction matter--dark energy were introduced by Wetterich \\cite{wetterich}. Meanwhile there exists a still growing body of literature on the subject -see, e.g. \\cite{interacting} and references therein. Limits for the admissible interaction strength have been obtained for various configurations \\cite{limits}. On the other hand, the transition from decelerated to accelerated expansion can be understood as a pure interaction phenomenon in the context of holographic DE models \\cite{DW,essay,WDCQG}.  Recently it was argued that the latest SNIa could favor a scenario in which the cosmic acceleration has past a maximum value and is now slowing down again \\cite{sastaro}. In such a case the accelerated expansion could be a transient phenomenon with a late-time dynamics, incompatible with that of the $\\Lambda$CDM model. It is this possibility which we are going to study in the present paper. We mention that models of transient acceleration were already discussed in \\cite{alcaniz} and more recently in \\cite{alcaniztr}.  It is the purpose of this paper to provide a toy model which describes a transient cosmological acceleration as the consequence of an interaction between dark matter and dark energy. Such a dynamics cannot be obtained if the interaction represents a small correction to, say, the $\\Lambda$CDM model. For models of this type  the long-time cosmological dynamics will always be determined by the cosmological term and result in accelerated expansion. To achieve transient accelerated expansion, a twofold role of the interaction is necessary. At first, it has to cancel the ``bare\" cosmological constant and at second it has to generate a phase of accelerated expansion by itself. Acceleration has to be an interaction phenomenon. We show that these requirements can be fulfilled by interaction terms that combine powers and exponentials of the cosmic scale factor.   The paper is organized as follows. In Sec. \\ref{Interacting fluid dynamics} we introduce the basic interacting fluid model and study the influence of the coupling, described by an initially arbitrary function of the scale factor, on the cosmic acceleration equation.  Sec. \\ref{transient} investigates two interaction types between DM and DE, each of them characterized by two parameters, in detail. Conditions for the existence of transient acceleration are found under the condition of a vanishing total effective cosmological constant. In Sec. \\ref{SN} we use a statistical analysis of the gold sample of the SNIa data \\cite{gold} to find the admissible parameter ranges for transient cosmic acceleration. The best-fit models for both interactions are shown to be compatible with the $\\Lambda$CDM model. Our conclusions are summarized in Sec. \\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}   In a homogeneous and isotropic cosmological background, a fluid with an equation of state $w=-1$ is dynamically equivalent to a cosmological constant. According to the $\\Lambda$CDM model, a small (in Planck units), but so far theoretically unexplained, cosmological constant is responsible for the observed accelerated expansion of our present Universe. We have demonstrated that an interaction between such type of dark energy and dark matter re-normalizes this constant. Even if the resulting effective cosmological constant is assumed to be zero, the time-dependent part of the interaction (the part that does not contribute to the re-normalization) was shown to be able to generate a phase of accelerated expansion of the Universe. We have studied two types of interactions for which cosmic acceleration is a transient phenomenon. While recent investigations favor a scenario in which the acceleration of the Universe is already slowing down today \\cite{sastaro}, our toy model predicts the maximal acceleration to occur at a future time. The $\\Lambda$CDM model is unable to provide any type of transient acceleration dynamics. A statistical analysis of the SNIa data shows that models of transient accelerated expansion may well be competitive with the standard cosmological constant model. While the mentioned cancelation mechanism of the ``bare\" cosmological constant is purely phenomenological, we believe that it might possibly indicate a more fundamental feature. A vanishing total effective cosmological constant seems more appealing and easier to understand within an underlying basic theory than a small, non-vanishing value of this quantity.  In general, our conclusion concerning the potentially transient character of the cosmic acceleration as the result of an interaction in the dark sector coincides with the numerical results of \\cite{alcaniztr}. But we think that an analytic solution, even though it is the solution of a toy model, provides additional information about the specific role of the interaction in this scenario and thus gives rise to a physically more transparent picture.         {\\bf Acknowledgement}  We thank FAPES and CNPq (Brazil) for financial support."
    },
    "0910/0910.1769_arXiv.txt": {
        "abstract": "\\n Quantum tunneling in Reissner--Nordstr\\\"om geometry is studied and the tunneling rate is determined. A possible scenario for cosmic inflation, followed by reheating phases and subsequent radiation-domination expansion, is proposed. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.0245_arXiv.txt": {
        "abstract": "Using samples drawn from the Sloan Digital Sky Survey, we study the relationship between local galaxy density and the properties of galaxies on the red sequence. After removing the mean dependence of average overdensity (or ``environment'') on color and luminosity, we find that there remains a strong residual trend between luminosity--weighted mean stellar age and environment, such that galaxies with older stellar populations favor regions of higher overdensity relative to galaxies of like color and luminosity (and hence of like stellar mass). Even when excluding galaxies with recent star--formation activity (i.e., younger mean stellar ages) from the sample, we still find a highly significant correlation between stellar age and environment at fixed stellar mass. This residual age--density relation provides direct evidence for an assembly bias on the red sequence such that galaxies in higher--density regions formed earlier than galaxies of similar mass in lower--density environments. We discuss these results in the context of the age--metallicity degeneracy and in comparison to previous studies at low and intermediate redshift. Finally, we consider the potential role of assembly bias in explaining recent results regarding the evolution of post--starburst (or post--quenching) galaxies and the environmental dependence of the type Ia supernova rate. ",
        "introduction": "\\label{sec_intro}  Over the past decade or more, the formation histories and clustering properties of dark matter halos, within a Lambda cold dark matter ($\\Lambda$CDM) cosmology, have been studied in detail using both numerical simulations and analytical analyses. These theoretical studies have repeatedly shown that more massive halos are more strongly clustered \\citep[e.g.,][]{gang88, cole89, mo96, sheth99, wetzel07}, a result that is supported by the observed clustering of nearby galaxy groups \\citep{padilla04, yang05, berlind06, wang08}. Furthermore, recent dark matter simulations indicate that the clustering of halos of a given mass depends on their assembly history, where halos that assembled their mass earlier in the history of the Universe are more strongly clustered than halos of the same mass that formed later \\citep{sheth04, gao05}; this effect is commonly referred to as assembly bias \\citep{croton07}.    The concept of assembly bias could be applicable to the formation of galaxies in addition to dark matter halos, such that galaxies that assembled their stellar mass earlier would be more strongly clustered today than galaxies of like mass and younger stellar populations. Such a bias in the galaxy population has typically been explored in two distinct manners: [1] by studying galaxy samples spanning a range of redshifts and directly observing the evolution in the population as a function of mass and clustering properties, and [2] by applying a more archaeological or paleontological approach, where the stellar populations of nearby galaxies are studied in detail, with the goal of disentangling the typical evolutionary history in different environments.   Various observational studies have investigated the formation of early--type or red--sequence galaxies from both perspectives. For instance, analyses of samples drawn from the DEEP2 Galaxy Redshift Survey \\citep{davis03, newman09} by \\citet{bundy06} and by \\citet{cooper06, cooper07} showed that the population of galaxies on the red sequence at $z < 1$ was preferentially built--up in overdense environments, so that at a given mass red galaxies in dense environs have typically ceased their star formation (i.e., assembled their mass) earlier than those in less dense regions. This picture is supported by analyses of the fundamental plane \\citep[FP,][]{dd87, dressler87} at similar redshifts, where it has been found that early--type galaxies in high--density regions have reached the FP more quickly than those in lower--density environs \\citep{vandokkum01, gebhardt03, treu05, moran05}.   In contrast to the relatively coherent picture derived from studies of early--type or red--sequence galaxies at intermediate redshift, detailed studies of the stellar populations in nearby ellipticals have yielded mixed results. A recent comparison of the stellar ages of early--type systems in the Coma cluster versus a corresponding field sample by \\citet{trager08} found no difference in assembly history between the two environment regimes. On the other hand, a similar analysis of elliptical and lenticular galaxies in the Virgo and Coma clusters by \\citet{thomas05} found a distinct relationship between mean stellar age and the environment in which a galaxy resides, such that cluster early--types are older than their field counterparts.    In this paper, we take a different approach towards studying the topic of assembly bias, with hopes of reconciling the results from studies at intermediate--redshift with those focusing on local samples. Many previous efforts at low redshift relied on comparing cluster and field populations and/or focused on high--quality observations of a relatively small number of galaxies and/or partially controlled for morphology while ignoring inter--sample variations in properties such as rest--frame color. In contrast, we utilize somewhat lower--quality data for a much larger sample of galaxies drawn from the Sloan Digital Sky Survey \\citep[SDSS,][]{york00} to study the assembly histories of thousands of red galaxies in a broad and continuous range of environments, controlling for correlations between environment and color, morphology, mass, etc. In the following Section (\\S \\ref{sec_data}), we describe the galaxy sample under study, including measurements of galaxy environments and stellar ages. Our results regarding the relationship between stellar mass, age, and environment are then presented in \\S \\ref{sec_results}, followed by further analysis, discussion, and a summary in \\S \\ref{sec_anal}, \\S \\ref{sec_disc}, and \\S \\ref{sec_summary}, respectively. Throughout, we employ a $\\Lambda$CMD cosmology with $w = -1$, $\\Omega_m = 0.3$, $\\Omega_{\\Lambda} = 0.7$, and a Hubble parameter of $H_0 = 100\\ h$ km s$^{-1}$ Mpc$^{-1}$, unless otherwise noted. ",
        "conclusions": "\\label{sec_disc}  In \\S \\ref{sec_results}, we presented the relationships between (both absolute and residual) environment and an assortment of galaxy properties, including $r$--band luminosity, rest--frame $g-r$ color, stellar mass, luminosity--weighted mean stellar age, and stellar metallicity. While we find no correlation between mean residual environment and stellar mass, we do observe significant relationships between $< \\! \\Delta_5 \\! >$ and stellar age and metallicity. We discuss the implications of these results in the remainder of this section.   \\subsection{Comparison to Previous Studies}  \\subsubsection{Age--Density Relation}  As highlighted in \\S \\ref{sec_intro}, studies of red--sequence galaxies in local clusters and within the field population have yielded discrepant results on the correlation between age and environment at fixed mass. On one hand, a variety of analyses studying early--type galaxies (i.e., E/S0 morphological types) in nearby clusters and the field have arrived at similar conclusions as our work, where red galaxies in higher--density regions are typically older at a given mass \\citep[e.g.,][]{trager00b, thomas05, clemens06, bernardi06, smith08}.\\footnote{However, some of these analyses have detected evidence of galaxy assembly bias at very low (or sometimes unspecified/undetermined) significance \\citep[e.g.,][]{thomas05, bernardi06, clemens06}.}  These studies utilized a variety of data sets from high--resolution spectroscopic observations of relatively small numbers of cluster members to composite spectra constructed from early SDSS data.   In contrast, more recent analyses have yielded contradictory conclusions. While \\citet{thomas05} found cluster ellipticals to have older stellar populations than their field counterparts at the same velocity dispersion or stellar mass, subsequent analysis by \\citet{thomas07} concluded that there is no significant environmental dependence to the slope or zero--point of the age--mass relation for the ``bulk'' of the nearby elliptical population (e.g., at ages of $\\gtrsim 4$ Gyr). Instead, \\citet{thomas07} claim that the trend detected by \\citet{thomas05} is driven entirely by a population of ``rejuvenated'' galaxies (i.e., ellipticals with recent star--formation activity and thus ages $\\lesssim \\! 3$ Gyr), which preferentially reside in low--density environs.   Our results disagree, in part, with the more--recent conclusions of \\citet{thomas07}. In agreement with their work, we find evidence for a variation in the fraction of ``rejunvenated'' galaxies with environment (that is, we find a residual age--density relation at young stellar ages). However, as shown in Figure \\ref{resid_fig} and Figure \\ref{select_fig}, even when restricting our sample to galaxies with ages of $\\log_{10}(t/{\\rm yr}) > 9.5$ (i.e., excluding the ``rejuvenated'' population), we detect a strong relationship between age and residual environment --- in conflict with the conclusions of \\citet{thomas07}. That is, we find that the correlation between luminosity--weighted mean stellar age and environment at fixed mass spans the entire range of ages probed by our various galaxy samples, rather than being an effect driven by only those galaxies with young stellar populations.   To emphasize and further illustrate this point, we select multiple subsamples according to various age selections (e.g., $t > 10^{9.5}$, $10^{9.6}$, $10^{9.7}$, and $10^{9.75}$ yr). For each of these samples of older galaxies, we independently repeat our analysis --- that is, we measure and remove the color--luminosity--density relation in each subsample and then examine the residual relationships between environment and color, luminosity, stellar mass, age, and metallicity. As shown in Figure \\ref{ages_fig} for a sample of red galaxies selected to have $\\log_{10}(t/{\\rm yr}) > 9.6$ and $r < 17.5$, we find no significant correlation between residual environment, $\\Delta_5$, and stellar mass. However, we do find significant residual age--density and metallicity--density relations, similar to those previously observed for samples selected according to apparent brightness and morphology yet spanning the entire range of ages probed (cf.\\ Figure \\ref{select_fig}).    \\begin{figure*}[tb] \\centering \\plotone{f7.eps} \\caption{The dependence of mean residual environment, $< \\! \\Delta_5 \\! >$, on absolute magnitude, rest--frame color, stellar mass, stellar age, and stellar metallicity for a sample of red--sequence galaxies with ages of $> \\! 10^{9.6}$ yr and $r < 17.5$. The points and error bars correspond to the mean environments and $1\\sigma$ uncertainties in the means computed in distinct bins of each galaxy property. While we find no significant residual trend between environment and color, luminosity, or stellar mass, we observe a strong dependence of residual environment on stellar age and metallicity.} \\vspace*{0.1in} \\label{ages_fig}  \\end{figure*}    The results as presented in Figure \\ref{ages_fig} are independent of the age limit employed to select the galaxy subsample; for samples with ages $> \\! 10^{9.5}$, $10^{9.7}$, and $10^{9.75}$ yr, we find no significant variation in the residual relationship with environment for stellar mass, age, and metallicity. Moreover, we find similar results for subsamples of galaxies selected according to a mass--dependent age criterion, analogous to that employed by \\citet{thomas07}. When removing the set of galaxies referred to as ``rejuvenated'' by \\citet{thomas07} (i.e., those with younger stellar populations), we still find a strong age--density relation at fixed stellar mass within the population of galaxies with older stellar populations (i.e., the ``bulk'' population according to \\citealt{thomas07}).   Finally, when defining subsamples selected on morphology and apparent brightness as well as age (e.g., $r < 17$, $n > 2.5$, and $\\log_{10}(t/{\\rm yr}) > 9.7$), we find trends between residual environment and stellar age and metallicity analogous to those presented in Figure \\ref{select_fig}. However, by excluding those galaxies with younger stellar populations in combination with morphological and apparent magnitude selections, the number of sources at low metallicities (e.g., $\\log_{10}(Z/Z_{\\sun}) < -0.5$) is significantly diminished such that the range in metallicities probed is limited and the resulting residual metallicity--density relation is only statistically significant within the metal--rich regime (i.e., $\\log_{10}(Z/Z_{\\sun}) \\gtrsim -0.35$). Still, altogether, this analysis illustrates that our results are inconsistent with the conclusions of \\citet{thomas07}, such that we find a relationship between stellar age and environment at fixed mass that spans all ages probed in our study (i.e., the age--density trend at fixed mass is not solely driven by galaxies with young stellar populations preferentially residing in underdense environs). Additional evidence for this conclusion, based on independent statistical tests, is presented in the Appendix.   Other recent analyses have, like \\citet{thomas07}, found little evidence of galaxy assembly bias at the massive end of the red sequence based on archaeological investigations of local galaxies. Studying a dozen early--type galaxies in the Coma cluster in detail, \\citet{trager08} also find no difference between the inferred stellar ages of red--sequence galaxies in the cluster environment and similar galaxies in the field. As discussed by \\citet{trager08}, the relatively young ages of the red galaxies in Coma are most likely attributed to recent star--formation episodes. However, nearly all of the galaxies in their sample show mean ages of $5$--$7$ Gyr, with none having stellar populations as old as $\\sim \\! 10$ Gyr. Past spectroscopic observations of early--type systems in Coma have generally arrived at a broader range of stellar ages, with a non--negligible population of galaxies with ages of roughly $10$ Gyr \\citep{moore01, poggianti01, moore02, nelan05, sanchez06}. However, the \\citet{trager08} data are of somewhat higher resolution and/or signal--to--noise ratio, facilitating improved emission corrections and precision in the resulting age estimates.    Accepting the accuracy of the \\citet{trager08} age measurements, our results are reconcilable if either the \\citet{trager08} sample is a biased tracer of the Coma early--type population (e.g., due to shot noise in the selection of their small sample) or if there are no (or at least very few) old early--type systems in the Coma cluster. If the latter is true, then Coma would be unusual amongst massive clusters in the local Universe. Coma is not a cool--core or cooling--flow cluster; instead it has a cooling time $\\gtrsim 10$ Gyr \\citep{kaastra04}. Thus, recent star formation at a significant level within the central region of the cluster seems unlikely, especially across the entire early--type population. However, Coma is not a relaxed system; instead, it contains significant substructure \\citep[e.g.,][]{biviano96, neumann03} and is observed to be undergoing a merger/accretion event \\citep[e.g.,][]{mellier88, neumann01}. As highlighted by \\citet{gerhard07}, a large fraction $(\\sim \\! 30\\%)$ of galaxies in Coma are likely associated with the accretion of the NGC 4839 subcluster. If Coma is a currently--forming supercluster, then perhaps recent episodes of star formation, associated with accretion/merger events, can explain the relatively young ages measured for the \\citet{trager08} sample of early--type galaxies.  All told, the young ages of early--type galaxies in Coma as found by \\citet{trager08} persists as a significant outlier relative to the results of other studies of local clusters, including other contemporary work studying Coma. For example, recent results from \\citet{smith09}, employing spectroscopic observations of comparable quality to those of \\citet{trager08} and spanning a larger sample of early--type galaxies extending to lower masses, find a significant population of galaxies with old stellar populations ($t \\sim 10$ Gyr) in the core of Coma. Moreover, \\citet{caldwell03} find massive early--type systems in Virgo with similarly old luminosity--weighted ages ($t > 9$ Gyr). While the sample of \\citet{smith09} is biased to lower masses than that of \\citet{trager08}, one would expect this to lead to smaller, not larger, ages than for the \\citet{trager08} sample, as more massive systems tend to contain older stellar populations \\citep[e.g.,][]{trager00a, trager00b, nelan05, rogers08, michielsen08, smith09, matkovic09}.   There are several other recent studies of local field and cluster galaxies that have found only weak statistical evidence for a relationship between stellar age and environment at fixed mass. For example, \\citet{gallazzi06} studied the environmental dependence of the scatter in the age--stellar mass relationship using the same catalog of stellar ages and metallicities as employed in this work. They find a relatively weak correlation between environment and age such that galaxies outside of the most dense environments tend to have younger stellar populations. The analysis of \\citet{gallazzi06}, however, was limited in its scope due to a rather small sample size ($< \\! 2000$ galaxies) for which environment estimates were available from \\citet{kauffmann04}. Our analysis expands upon this earlier work by markedly increasing the sample size with environment information and by utilizing a more sensitive statistical methodology.   While \\citet{clemens06} and \\citet{bernardi06} both conclude that there are differences between the ages of massive ellipticals in differing environments locally, with systems in high--density regions being roughly $1$ Gyr older at fixed velocity dispersion, evidence for a significant age--density relation at fixed mass in their samples is considerably weaker than that presented herein and more in line with the results of \\citet{thomas07}. For example, \\citet{clemens06} find a significant difference between the ages of ellipticals in high-- and low--density environs over only a limited range in velocity dispersion (or specifically for only 2 of their 7 discrete bins in velocity dispersion), with the differences (in those two bins) significant at less than $2\\sigma$ \\citep[see Figure 10 of][]{clemens06}. Similarly, in the work by \\citet{bernardi06}, the apparent difference in average age between ellipticals in high--density and low--density environs depends on the particular Balmer absorption line (H$\\beta$ versus H$\\gamma_{\\rm F}$) employed to constrain stellar age, with ages based on H$\\beta$ showing no significant difference as a function of environment and with H$\\gamma_{\\rm F}$ yielding age variations significant at the approximately $2\\sigma$ level \\citep[see Table 9 of][]{bernardi06}.  The lower significance at which \\citet{bernardi06} and \\citet{clemens06} find a trend between stellar age and environment at fixed mass likely results from the smaller sample sizes, alternate methods of measuring environment, and different statistical methods employed in the analysis. For example, by selecting samples drawn from the extremes of their measured environment distributions, both \\citet{bernardi06} and \\citet{clemens06} are limited to final sample sizes of $< \\! 9000$ and $< \\! 4000$ galaxies, respectively. On top of that, \\citet{clemens06} further reduce the statistical power of their sample by dividing into course environment bins and then subdividing into bins of velocity dispersion \\citep[e.g., see Figure 10 of][]{clemens06}. While \\citet{clemens06} uses an environment measure similar to ours, \\citet{bernardi06} use the 3--dimensional distance to the $10^{\\rm th}$--nearest neighbor with $M_r < -21.5$ as an environment estimator. By restricting the tracer population to relatively luminous galaxies, while also using a $10^{\\rm th}$-- versus $5^{\\rm th}$--nearest--neighbor distance (as used in this work), \\citet{bernardi06} are tracing environment on much larger scales, perhaps missing the dependence of galaxy properties on the more local galaxy density. Moreover, measuring distances in 3--dimensions versus in projection leads to greater sensitivity to the effects of redshift--space distortions, which smear out the galaxy distribution in environments such as clusters \\citep{cooper05}. Finally, along with using the distance to the $10^{\\rm th}$--nearest neighbor, \\citet{bernardi06} use the distance to the nearest galaxy cluster as a means for identifying high--density environments, thereby significantly biasing their sample against group--like environs.  Altogether, these effects likely lead to the marginally--significant trends between stellar age and environment at fixed mass presented by \\citet{bernardi06} and \\citet{clemens06}. In contrast, our methodology, which utilizes the entire galaxy sample (versus only studying galaxies in extreme environs), yields a strong age--density relation at fixed mass.   \\subsubsection{Metallicity--Density Relation}  In addition to a strong residual age--density relation, we also find a relatively weak residual metallicity--density relation, where more metal--rich galaxies favor regions of higher galaxy density relative to galaxies of like stellar mass. However, the significance of this metallicity--density relation decreases as the sample is restricted to more centrally--concentrated (i.e., more bulge--dominated) or less elongated galaxies (cf.\\ \\S \\ref{sec_disc1}), such that a weak (though statistically significant) correlation is observed between mean residual environment, $< \\!  \\Delta_5 \\! >$, and stellar metallicity for the bulk of the galaxy samples studied here.   Previous analyses of early--type galaxies within local clusters have generally found no evidence for a metallicity--density relation at fixed mass, in contrast to our results. For instance, studies of line indices in nearby clusters \\citep[e.g.,][]{guzman92, carter02} find a significant dependence of the Mg$_2$--$\\sigma$ relation on cluster--centric distance, where early--type galaxies near the core exhibit stronger Mg$_2$ absorption than their counterparts in the outskirts of the cluster. As \\citet{guzman92} only measure one spectral index, they are unable to disentangle age and metallicity; although they do conclude that a dependence of stellar age on cluster radius might be responsible for the observed variations in Mg$_2$ absorption strength \\citep[see also][]{bernardi98}. A similar conclusion is reached by \\citet{terlevich01} studying the $U-V$ color of galaxies in the Coma cluster; color, however, is similarly unable to distinguish variations in age from variations in metallicity.   Folding in additional spectral indices to separate metallicity and age effects, \\citet[][see also \\citealt{kuntschner98}]{carter02} infer a significant relationship between metallicity and radius in Coma. However, this metallicity--density relation does not take into account the relationship between mass (or luminosity) and radius, minimizing its usefulness in a comparison to our analysis. More recent work by \\citet{smith06}, measuring a larger number of indices ($12$) for a significantly larger sample of galaxies ($\\sim \\! 3000$) in $94$ nearby clusters, finds several significant relationships between spectral indices and cluster--centric radius at a given velocity dispersion. Using these indices to distinguish between age and metallicity effects, they conclude that there is a significant age gradient in local clusters, such that systems in cluster cores have older stellar populations, while there is no significant metallicity gradient detected \\citep[see also][]{spolaor09}.    Other recent analyses spanning a broad range of environments in the local Universe have also found little evidence for a significant relationship between stellar metallicity and local galaxy density at fixed mass. For example, both \\citet[][see also \\citealt{clemens09}]{clemens06} and \\citet{bernardi06}, find no significant metallicity--density trend at fixed velocity dispersion for nearby elliptical and lenticular systems \\citep[see also][]{gallazzi06}. Furthermore, while \\citet{thomas05} conclude that massive early--type galaxies in local clusters are, on average, $\\sim \\! 0.05$--$0.1$ dex more metal--poor than their counterparts in the field, the significance of this result is unsubstantiated statistically, with no uncertainties given for the coefficients of the linear fits to the mean $Z$--$\\sigma$ relations in the high-- and low--density galaxy samples \\citep[see Equation 1 of][]{thomas05}. Considering the small size of the galaxy sample (only $124$ galaxies in total) employed by \\citet{thomas05}, the uncertainties on the parameters of the linear fits is likely to exceed a few percent, which would make the small measured differences in the mean $Z$--$\\sigma$ relations in the two environment regimes statistically insignificant. Thus, \\citet{thomas05}, like many other authors, find no evidence for a metallicity--density relation within the nearby early--type galaxy population.   Not only do we find a significant metallicity--density relation at fixed mass for our most inclusive color--selected sample, but also for our morphologically--selected subsamples. While this trend is weaker than the age--density relation at fixed mass, it is significant nonetheless. As discussed in the preceding section (\\S \\ref{sec_degen}), a metallicity--density relation is unlikely to result from correlated errors associated with the degeneracy between constraints on stellar age and metallicity. Moreover, when comparing apples to apples, the uncertainties in measured galaxy properties are generally large enough that it is far easier to dilute an underlying environment trend, such that it goes undetected, than it is to falsely detect one. Thus, for many previous studies of local early--type galaxies, it is likely that the same errors or limitations that smeared out the age--density relation in the galaxy sample also precipitated a non--detection of the metallicity--density relation.   In contrast to our results and to the majority of results from the aforementioned studies of local early--type galaxies, \\citet[][see also \\citealt{rose94}]{kuntschner02} find that their sample of nearby field ellipticals are more metal--rich (by $\\sim \\! 0.2$ dex) in comparison to members of the Fornax cluster. This discrepancy between the findings of \\citet{kuntschner02} and those of our work (as well as those of \\citet{bernardi06}, \\citet{clemens06}, among others) can potentially be understood in terms of an age--metallicity degeneracy. As discussed in more detail in \\S \\ref{sec_degen}, errors in estimating stellar age are strongly correlated with errors in metallicity, such that if a galaxy's age is overestimated, then its metallicity will correspondingly be underestimated. Thus, the metallicity--density relation reported by \\citet{kuntschner02} could be reconciled with other observations, if the \\citet{kuntschner02} age--density relation is correspondingly stronger than that found in other analyses. The plausibility of this scenario is supported by the larger difference in age between massive ellipticals in field and cluster environments given by \\citet[][$\\sim \\! 2$--$3$ Gyr]{kuntschner02} versus that found by \\citet[][$\\sim \\! 1$ Gyr]{bernardi06}.\\footnote{Note that Table 9 of \\citet{bernardi06} lists an age difference between high-- and low--density environs that is more in line with $\\sim \\! 0.6$ Gyr.}   The analysis by \\citet{kuntschner02} suffers even more from the problem of small sample size ($< \\! 30$ galaxies in total). In addition, there was a strong difference in the morphological composition of their field and cluster samples, with the field population containing a larger percentage of S0s (versus ellipticals). However, the \\citet{kuntschner02} study compared the ages of galaxies in different environments at fixed $B$--band absolute magnitude and not at fixed mass (i.e., not at fixed stellar mass or velocity dispersion). Thus, their field sample may have been increasingly dominated by younger galaxies at lower masses, which appeared bright in the $B$--band due to recent star--formation activity. Given the observed correlation between mass and metallicity (independent of environment), such a bias in the sample would work to decrease the strength of the metallicity--density relation detected by \\citet{kuntschner02}.     Recognizing the correlation between stellar metallicity and age observed for local massive early--type systems, where galaxies of a given mass with older stellar populations tend be more metal--poor \\citep[e.g.,][]{gallazzi05, sanchez06}, it might be expected that an age--density relation at fixed mass would naturally be accompanied by some sort of metallicity--density relation, like that observed by \\citet{kuntschner02}. However, there is large scatter in the age--metallicity relation, especially at lower masses. Furthermore, stellar age and metallicity measure different aspects of a galaxy's formation history; while age traces the time since the bulk of the stars were formed (or in the case of luminosity--weighted mean stellar age, the time since the last significant episode of star formation), metallicity reflects the balance of feedback and accretion processes in concert with the star--formation efficiency (that is, the efficiency with which gas is turned into stars). For example, two galaxies forming from pristine gas reservoirs and following identical formation histories, except for an offset in time, will contain stellar populations of comparable metallicity, but distinct ages. Our results suggest that not only did red--sequence galaxies in high--density environments form earlier than their counterparts in low--density regions, but [1] feedback and/or accretion processes are modulated by the local environment such that the high--density regime favors greater metal abundances and/or [2] the typical star--formation efficiency of galaxies in high--density regions or at higher redshift exceeds that of galaxies in low--density environs or at lower redshift.  Some evidence for the former scenario has been observed in a study of star--forming galaxies at $z \\sim 0.1$ by \\citet{cooper08b}. Using the gas--phase oxygen abundance measurements of \\citet{tremonti04}, they show that local environment is correlated with a non--negligible portion of the scatter in the mass--metallicity relation, such that galaxies in higher--density environs at a given mass tend to be more metal--rich. As discussed in more detail by \\citet{cooper08b}, there are many physical scenarios that might lead to galaxies having a higher metal abundance in high--density regions versus in low--density regions. For example, processes such as ram--pressure stripping or galaxy harassment \\citep{gunn72, moore96}, which only operate in extreme high--density environs, could remove the outer portion (i.e., the most metal--poor portion) of a galaxy's gaseous halo, thereby inflating the relative metal abundance of gas reservoirs within galaxies in rich groups and clusters. Similarly, outflows and winds driven from galaxies in high--density environments might be less capable of expelling enriched gas versus their counterparts in underdense regions.  There is also evidence to support the possible higher star--formation efficiencies of star--forming galaxies at higher redshift relative to local star-forming galaxies. For example, studies of sub--millimeter galaxies \\citep[SMGs,][and references therein]{chapman05} as well as luminous and ultra--luminous infrared galaxies \\citep[LIRGs and ULIRGs,][]{aaronson84, sanders96} at $z > 1$ typically find higher surface densities of star formation for a given surface density of molecular gas \\citep[][]{bouche07, daddi08, tacconi09}. Such a correlation between star--formation efficiency and redshift in concert with a galaxy assembly bias where galaxies in high--density environs formed their stars earlier (i.e., at higher redshift) than their counterparts in low--density environs would yield a corresponding correlation between environment and star--formation efficiency such that galaxies in overdense regions are more efficient at converting gas into stars (i.e., have higher stellar metallicity).      \\subsection{The Impact of Mergers}  The stellar populations of galaxies are not entirely formed via in--situ star formation. Rather, much of a galaxy's stellar mass is assembled via accretion events (i.e., mergers). In addition to increasing the stellar mass of a galaxy, mergers can also have a substantial impact on the age of the resulting system. While gas--rich (or ``wet'') mergers will tend to induce bursts of star formation thereby significantly decreasing the inferred luminosity--weighted mean stellar age of the system, dissipationless (or ``dry'') merging of two systems with older stellar populations will have a less severe impact on the inferred age of the merger product, yielding a passively--evolving descendant with a stellar age that is simply the luminosity--weighted average of the antecedents' ages.   In addition, mergers tend to occur in regions of higher galaxy density, such as galaxy groups \\citep[e.g.,][]{cavaliere92, mcintosh08, hester09, ideue09, lin09}. Given this correlation with environment, it is important to understand the potential impact of hierarchical assembly on our results. Since mergers increase the stellar mass of the system and since galaxies with smaller stellar masses tend to have younger stellar populations \\citep[e.g.,][]{gallazzi05, graves09}, galaxies that have recently undergone a merger are likely to be biased towards lower ages relative to other galaxies of like stellar mass. In the case of gas--rich mergers, this effect is further accentuated by the decrease in age associated with the merger--induced burst of star formation. Thus, mergers would tend to decrease the average age of galaxies at a given stellar mass in the high--density regime, thereby decreasing the amplitude of the observed age--density relation at fixed mass. In all, the impact of mergers would be to smear out the residual age--density relation shown in Figure \\ref{resid_fig}.     \\subsection{The Evolution of Post--Starburst or Post--Quenching Galaxies}  Our analysis of red galaxies in the local Universe shows a clear relationship between stellar age and environment, where galaxies with older stellar populations tend to reside in higher--density environments relative to younger galaxies of like stellar mass. Such a correlation between stellar age and environment could naturally explain the evolution in the clustering of post--starburst (otherwise known as K+A or post--quenching) galaxies at $z < 1$, as observed by \\citet{yan09}.  Using data drawn from the SDSS and from DEEP2, \\citet{yan09} studied the distribution of environments for K+A galaxies \\citep{dressler83} at $z \\sim 0.1$ and at $z \\sim 0.8$, relative to the red and blue galaxy populations at those redshifts. They find that post--starburst galaxies at low redshift have an environment distribution similar to that of blue galaxies, favoring regions less dense than those typically inhabited by red galaxies \\citep[see also][]{hogg06}; however, at higher redshift, they find post--starburst systems favor environments more similar to that of red galaxies.  If post--starburst (or K+A) galaxies have their star--formation quenched in the same type of environment (e.g., a density level corresponding to a galaxy group) at all redshifts less than $z \\sim 1$, then this evolution in the environment distribution of post--starbursts relative to that of the red galaxy population from $z \\sim 0.8$ to $z \\sim 0.1$ can naturally be explained in terms of galaxy assembly bias. On average, a red galaxy at low redshift will have had a longer amount of time since having its star formation quenched; thus, the environment of a low--$z$ red galaxy will have had more time to increase in density relative to that in which the quenching occurred (i.e., the environment in which the K+A phase of evolution occurs).  In this scenario, the time since quenching (or time since assembly --- thus, stellar age) for a galaxy on the red sequence is naturally correlated with the local galaxy density or environment. Our results support this potential explanation for the evolution of the environment distribution of post--starburst galaxies at $z < 1$. The observed correlation between stellar age and environment at fixed stellar mass supports (although does not prove) the picture in which K+A galaxies are found in the same type of environment at $z < 1$.     \\subsection{Implications for the Environmental Dependence of the Type Ia SN Rate}  The typical age of a stellar population has been shown to be strongly connected to the type Ia supernova (SN) rate, with young star--forming galaxies having much higher SN Ia rates than quiescent galaxies with older stellar populations \\citep{sullivan06}. The increase in the SN Ia rate with star--formation rate is commonly attributed to the existence of two types of Ia events: ``prompt'' SNe Ia, which likely result from the evolution of somewhat massive stars and are thus correlated with star formation, and ``delayed'' SNe Ia, which are thought to be the evolutionary outcome of less massive stars and therefore connected to the underlying stellar mass of the galaxy.  Within quiescent systems, the delayed component of the Ia rate dominates. However, even for systems of like mass and lacking ongoing star formation, the SN Ia rate is still expected to depend on the age of the stellar population; some theoretical predictions for both single--degenerate and double--degenerate progenitor scenarios suggest that the SN Ia rate for a single--aged stellar population should increase with time, reaching a peak around an age of $\\sim \\! 10^{8}$ yr, with a sharp decline towards later times \\citep{greggio83, yungelson94}. A more recent analysis by \\citet{greggio05} slightly revises this general picture, finding a sharp reduction in the type Ia rate for both single--degenerate and double--degenerate progenitors beyond an age of $\\sim \\! 1$ Gyr.   Taking as given this theoretical connection between stellar age and type Ia rate, our results regarding the correlation between age and environment at fixed stellar mass along the red sequence suggest that the type Ia rate in early--type galaxies should vary with environment. In particular, our study indicates that the SN Ia rate within quiescent early--type galaxies should be lower in high--density environs relative to that found in low--density environs. The bulk of our samples have typical luminosity--weighted mean stellar ages greater than $1$ Gyr, placing these populations beyond the expected peak in the type Ia rate and into a regime where the rate should be in sharp decline.   Currently, the evidence is mixed regarding a potential correlation between the SN Ia rate (per unit mass) and environment in the local Universe. Most notably, \\citet{mannucci08} concluded that the type Ia rate is more than 3 times higher in cluster ellipticals relative to field ellipticals, using a sample of 11 cluster and 5 field type Ia events at $z < 0.04$, which would run counter to the theoretical trend. However, the uncertainty of these environment--dependent rates is quite large. In contrast, a recent analysis of a larger sample of type Ia supernovae drawn from the SDSS--II Supernova Survey by \\citet{cooper09} found no significant relationship between local galaxy environment and the likelihood of hosting a type Ia event within the red galaxy population. Overall, such studies of the environments of type Ia supernovae are currently limited by small numbers. Results from future surveys \\citep[e.g.,][]{sharon07, sand08} will hopefully help in answering such outstanding questions.     Using the measurements of luminosity--weighted mean stellar age and stellar metallicity from \\citet{gallazzi05} and the local galaxy overdensity estimates of \\citet{cooper09}, we study the relationship between galaxy properties and environment for various samples of galaxies on the red sequence. In contrast to previous studies in the local Universe, we do not solely focus on the early--type galaxy population or on galaxies selected to reside in particular environments (e.g., specific clusters such as Virgo or Coma); instead, we define various galaxy samples drawn from the SDSS, controlling for correlations between environment and color, morphology, mass, etc., while probing a broad and continuous range of environments. Our principal results are as follows:  \\begin{enumerate}  \\item After removing the mean color--luminosity--environment relation on the red sequence, we find a strong residual relationship between environment and stellar age (cf.\\ Figure \\ref{resid_fig}), such that galaxies with older stellar populations favor regions of higher galaxy overdensity relative to galaxies of like color and luminosity (i.e., like stellar mass).  \\item When using subsamples restricted according to morphological, brightness, and color criteria, we find this residual age--density relation persists with no significant change in strength (cf.\\ Figure \\ref{select_fig} and Table \\ref{tab2}).  \\item Using stellar velocity dispersion to trace mass (in lieu of stellar mass), we find that the age--density relation at fixed mass persists (cf.\\ Figure \\ref{vdisp_fig}) and is even more statistically significant (cf.\\ Table \\ref{tab1} and Table \\ref{tab3}).  \\item We find a significant residual correlation between environment and stellar metallicity at fixed color and luminosity (cf.\\ Figure \\ref{resid_fig}). While a strong residual $Z$--density relation is observed for our most inclusive sample of red--sequence galaxies, enforcing selection criteria based on concentration, axis ratio, and/or brightness yields samples that exhibit weaker (though still statistically significant) residual metallicity--density relations (cf.\\ Figure \\ref{select_fig} and Table \\ref{tab2}).  \\item Our results are in conflict with the general conclusions of recent work by \\citet{thomas07} and \\citet{trager08}, which both indicate a lack of correlation between age and environment at fixed mass. While the results of \\citet{trager08} may be due to the peculiarities of the Coma cluster (relative to other local clusters) and/or sample selection effects, the \\citet{thomas07} findings are clearly not supported by our analyses. While we find evidence for a variation in the fraction of ``rejuvenated'' (or young) galaxies with environment \\citep[in agreement with the results of][]{thomas07}, we also find a correlation between age and environment at fixed mass even for the ``bulk'' (or old) early-type galaxy population --- counter to the claims of \\citet{thomas07}.   \\item In agreement with several previous studies of early--type systems in the local Universe and with a multitude of studies at intermediate redshift, we find significant evidence for assembly bias on the red sequence, such that red, early--type galaxies that formed earlier in the history of the Universe are more strongly clustered today than their counterparts (of equal mass) that formed later.   \\end{enumerate}   As discussed in \\S \\ref{sec_intro}, evidence for an assembly bias in the galaxy formation process has commonly been explored in two distinct manners: an evolutionary approach, which studies galaxies over a range of redshift and directly infers evolution in the galaxy properties, and an archaeological or paleontological approach, which studies the stellar populations of nearby galaxies in detail with the aim of deriving the evolutionary history from the current fossil record. As evidenced by the work presented herein, the results from these two methods are now arriving at a relatively consistent picture. Studies of galaxies at $z < 1$ have shown that the population of red (or early--type) galaxies was preferentially built--up in overdense environments, such that red--sequence galaxies in dense environments have generally assembled their stellar mass and developed elliptical morphologies earlier than those in less dense regions \\citep[e.g.,][]{treu05, bundy06, cooper07, gerke07, capak07}. In addition, studies of galaxy populations at yet higher redshifts ($z \\sim 2$--$3$) have detected populations of massive galaxies with passively--evolving stellar populations \\citep[e.g.,][]{labbe05, longhetti05, papovich06, stutz08} that are very strongly clustered \\citep[e.g.,][]{daddi03, quadri07, quadri08}.  Our results, which apply the aforementioned archaeological method, show a general agreement with these results from studies at higher redshifts, supporting a picture of assembly bias in the local galaxy population. We find that at a given mass on the red sequence galaxies with older stellar populations are found in more overdense environments (i.e., are more strongly clustered) than their counterparts with younger stellar populations. In contrast to some previous studies, we find that this correlation between age and environment at fixed mass is \\emph{not} solely driven by a minority population of galaxies that have experienced recent episodes of star formation. Instead, we find a strong residual age--density relation even among those red--sequence galaxies dominated by older stellar populations (i.e., with ages $\\gtrsim \\! 5$ Gyr).  In the future, we look to quantify the variation in the age of stellar populations as a function of environment, with the goal of constraining theoretical models of galaxy formation and evolution. As shown in the Appendix, preliminary analysis of only those galaxies with older stellar populations (i.e., $\\log_{10}(t/{\\rm yr}) > 9.75$) clearly demonstrates that there is a difference of \\emph{at least} 0.25 Gyr between those systems residing in high--density environs and those in low--density regions. Some recent analyses at intermediate redshift \\citep[e.g.,][]{vandokkum03, vdw05, vandokkum07} suggest that while early--type galaxies in clusters are older than their field counterparts, the difference in stellar age is less than would be expected from current generations of galaxy formation models. More quantitative analysis at low and intermediate redshift will be required to definitively measure the relationship between stellar population parameters and environment at $z < 1$ and thereby further constrain the theoretical models.     \\vspace*{0.25in}"
    },
    "0910/0910.0842_arXiv.txt": {
        "abstract": "In a recent publication, an aperture mass statistic for gravitational flexion was derived and shown to be effective, at least with simulated data, in detecting massive structures and substructures within clusters of galaxies. Further, it was suggested that the radius at which the flexion aperture mass signal falls to zero might allow for estimation of the mass or density profile of the structures detected. In this paper, we more fully explore this possibility, considering the behaviour both of the peak signal and the zero-signal contours for two mass models--the singular isothermal sphere and Navarro-Frenk-White profiles--under varying aperture size, filter shape and mass concentration parameter. We demonstrate the effectiveness of the flexion aperture mass statistic in discriminating between mass profiles and concentration parameters, and in providing an accurate estimate of the mass of the lens, to within a factor of $1.5$ or better. In addition, we compare the aperture mass method to a direct nonparametric reconstruction of the convergence from flexion measurements. We demonstrate that the aperture mass technique is much better able to constrain the shape of the central density profile, obtains much finer angular resolution in reconstructions, and does not suffer from ambiguity in the normalisation of the signal, in contrast to the direct method. ",
        "introduction": "There is an ongoing debate regarding the exact shape of galaxy- and cluster-scale dark matter haloes in the Universe. N-body simulations carried out under the assumption of the standard concordance cosmological model, $\\Lambda$CDM, suggest that cold dark matter haloes are well fit by Navarro-Frenk-White (NFW) density profiles (Navarro, Frenk \\& White, 1997), which are defined by a characteristic virial radius, $R_{200}$ (or, equivalently, the virial mass $M_{200}$), and a concentration parameter, $c$. However, gravitational lensing observations of field galaxies find that galaxy haloes consistently appear to show isothermal density profiles on the scales probed by strong lensing (Treu \\& Koopmans, 2002; Rusin, Kochanek \\& Keeton, 2003; Rusin and Kochanek, 2005; Koopmans et al., 2006; Gavazzi et al., 2007; Czoske et al., 2008; Dye et al., 2008; Tu et al., 2009), resulting from the interplay between baryons and dark matter. There is further debate concerning the density profiles of galaxy clusters (Carlberg et al., 1997; van der Marel et al., 2000; Athreya et al., 2002; Katgert, Biviano \\& Mazure, 2004; Lin, Mohr \\& Stanford, 2004; Hansen et al., 2005; {\\L}okas et al., 2006; Rines \\& Diaferio, 2006; Wojtak et al., 2007; Okabe et al., 2009); although some gravitational lensing studies favour an NFW model over an isothermal model, there are claims that NFW models give a poor fit to some clusters (e.g. Newman et al., 2009).  In addition, there has been much discussion regarding the relationship between halo mass and NFW concentration parameter. $\\Lambda$CDM N-body simulations imply that the NFW profile is, in fact, a one-parameter (rather than two-parameter) model with the concentration parameter, $c$, being related directly to the virial mass $M_{200}$ as a power law with negative index (e.g. Navarro et al., 1997; Shaw et al., 2006). Gravitational lensing observations involving galaxy-galaxy lensing in the Sloan Digital Sky Survey, for example, seem to support this supposed mass-concentration relation (Mandelbaum, Seljak, \\& Hirata, 2008). However, numerous gravitational lensing studies of clusters of galaxies, both in the strong and weak lensing regimes, have found significantly higher concentration parameters than would be expected from the mass-concentration relation (e.g. Abell 1689, for which a compilation of results can be found in Corless, King \\& Clowe, 2009).  Furthermore, in a number of gravitational lensing studies of clusters of galaxies, different sets of lensing measurements are seen to disagree, and different (NFW) parametric models of the cluster with widely varying masses and concentration parameters are seen by different authors to provide good fits to the data. For example, gravitational lensing studies of the galaxy cluster Abell 1689 (see, e.g., Corless et al. 2009; Peng et al. 2009; and references therein) have found best-fit models that span a factor of 3.5 in mass and 3 in best-fit NFW concentration parameter. Such disagreement hinders efforts to compare cluster data with the predictions from N-body simulations described above, and thus to constrain cosmological parameters. The technique presented in this paper is designed to assist in breaking some of these degeneracies, as even within a family of models such as the NFW model, the $\\fmap$ signatures for profiles with different lens masses and concentration parameters are seen to be substantially different from one another.   Measurements of gravitational flexion, the second order gravitational lensing effects which give rise to skewness and arciness in galaxy images, have been shown by numerous authors to be adept at detecting galaxy- and galaxy group-size haloes both in the field and within clusters of galaxies (e.g. Okura, Umetsu \\& Futamase, 2007,2008; Leonard et al., 2008; Leonard, King \\& Wilkins, 2009), as well as providing an alternative method by which the mass distribution within a cluster of galaxies might be constrained. In addition, Lasky \\& Fluke (2009) have found that the first flexion signal from an SIS profile differs from that of an NFW substantially at moderate separation between source and lens, whilst the second flexion signal shows strong variation when the NFW concentration parameter is varied. Flexion studies could therefore provide complementary constraints on halo profiles as well as their masses.  Recently, Leonard et al. (2009; hereafter LKW09) developed an aperture mass statistic for flexion, in direct analogy with that used for shear, and showed that it provided a robust method by which structures within clusters of galaxies, on a range of physical scales, can be identified to appreciable signal to noise. This technique is formally identical to the standard shear aperture mass methods used in weak lensing (see, for example, Schneider 1996, Schneider et al. 1998), but uses measurements of flexion rather than shear. This technique has the advantages that the noise properties of the aperture mass maps generated are very well understood, and that the filter functions used can be tuned to provide optimal signal-to-noise for a given lens profile.  Moreover, the flexion aperture mass technique offers a robust method by which the mass distribution of a cluster of galaxies might be mapped out without the need for parametric modelling. This means that the technique does not rely on assumptions about the mass profiles of the structures responsible for the lensing signal, nor does it require \\textit{a priori} knowledge of the locations of the mass concentrations responsible for the lensing signal. It can therefore be used to place independent constraints on the mass distribution of clusters of galaxies, without the invocation of any assumptions regarding the shape of the mass density profile of the structure.  In addition, LKW09 noted that the zero-signal contours expected to be found around mass peaks vary in radius under change of mass profile, lens mass, filter shape and aperture size. LKW09 suggest the properties of this set of contours might thus be used to provide insights into the mass and profile shape of the structures identified using this method, thus offering the potential to discriminate between competing mass models in cases where degeneracies arise. In this paper, we more fully investigate this claim.  We focus on two properties of the flexion aperture mass signal--namely, the peak signal associated with a given structure and its zero-signal contour--and investigate their behaviour for SIS and NFW profiles under variation in virial mass, NFW concentration parameter, aperture radius and filter shape in order to quantify the discriminating power of this technique. We demonstrate that the flexion aperture mass signatures of these two mass profiles are very different under changes to the aperture and filter properties. As a result of this divergent behaviour, aperture mass filtering of flexion data using a variety of aperture parameters provides a convenient and straightforward method for estimating both the mass and profile shape of the structures detected without the need for parametric modelling. Moreover, we demonstrate on simulated data that the total mass can be constrained to within a factor of $\\sim 1.5$ for both galaxy-group and cluster scale haloes, and the input profile shape recovered, for a signal to noise in the flexion aperture mass measurement as low as 1. Finally, we demonstrate that this method significantly outperforms a Fourier transform-based direct inversion technique (similar to that used in Okura et al.'s 2008 analysis of Abell 1689), yielding both higher resolution and much better constraints on the shape of the cluster mass profile.  This paper is structured as follows. In \\S~\\ref{sec:formalism}, we provide a brief review of the origin of the flexion signal, an overview of the flexion aperture mass statistic, and a description of the mass models considered. In \\S~\\ref{sec:radprof}, we consider the radial profile of the aperture mass  for each of the mass models and for a range of filters and aperture radii, and for varying values of the lens virial mass and concentration parameter (in the case of NFW lenses). The behaviour of the peak signal and the zero-signal contour under changes to the underlying mass model and changes to the aperture and filter properties is considered. In \\S~\\ref{sec:discrim}, we investigate the behaviour of a single mass profile under changes to the aperture and filter properties, and examine whether the use of several different combinations of aperture radius and filter shape in aperture mass reconstructions enables one to discriminate between different halo masses and density profiles. In \\S~\\ref{sec:nfw}, we apply aperture mass and direct reconstruction techniques to flexion data obtained by raytracing simulations through a cluster extracted from the Millennium simulation (see LKW09 and references therein) and compare the performance of the two methods. We conclude in \\S~\\ref{sec:summary} with a discussion of our results and their implications for future flexion studies.  Throughout the text, we assume a standard $\\Lambda$CDM cosmology with $\\om=0.27$, $\\Omega_\\Lambda=0.73$, and $H_0=100h$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$. ",
        "conclusions": "\\label{sec:summary}  In this paper, we have presented the expected $\\fmap$ signal from two common mass profiles: the Singular Isothermal Sphere model and the Navarro-Frenk-White model. This signal is characterised by a peak value, found when the aperture location coincides with a mass concentration, and a zero-signal contour. By simulating radial $\\fmap$ profiles for a large number of combinations of halo mass, NFW concentration parameter, aperture radius and filter polynomial order, the behaviour of the peak signal and zero-signal contour has been characterised for both the SIS and NFW density profiles under the assumption of circular symmetry.  The behaviour of both these measures is seen to diverge rather sharply between different models with the same mass as the aperture properties are varied. This implies that one might be able to use the $\\fmap$-statistic to constrain the masses and mass profiles of structures detected simply by varying the aperture parameters and noting the change in the peak signal and the location of the zero-signal contour.  Indeed, it has been shown using simulated data from analytic models that with modest signal to noise and relatively large errors on the measurement of the zero-signal radius, it is possible to isolate an input model using a reasonable number of combinations of aperture size and filter polynomial order. It also becomes easier to discriminate between masses at the high-mass end of the spectrum considered here, with higher mass haloes requiring fewer combinations of $l$ and $R$ to reach convergence.  This was further reinforced by testing the method on $\\fmap$ data using the simulated cluster of LKW09 as the input lens. This lens has a virial mass of $1.23\\times10^{15}h^{-1}\\,M_\\odot$, and the combination of three $\\fmap$ measurements of the lens was able to recover this mass to within a factor of $<1.5$, and to correctly identify the shape of the profile as an NFW with low concentration parameter. This success, despite the circularisation of a clearly elliptical lens profile, demonstrates solidly the power of the $\\fmap$-statistic to discriminate between haloes of different masses.  Moreover, this method shows a dramatic improvement over direct nonparametric reconstruction techniques. Firstly, the resolution that one can obtain using a direct method is limited to a far greater extent than aperture methods. This is because the number of bins one can use is limited by the density and distribution of background sources. We find, for the simulated cluster, a best resolution of just under half an arcminute. This coarse resolution is not problematic with the simulated cluster presented here, as it has a fairly smooth large-scale distribution. However, where there are significant substructures (such as in Abell 1689), averaging a signal over large bins has the effect of washing out the signal from smaller structures.  Aperture measures, on the other hand, filter the signal through a large aperture ($R>0.5$ arcminutes, in general), but are able to obtain fine resolution in the output aperture mass by spacing the apertures close together. Thus, finer structure can be resolved using these methods.  In addition, we have shown that the ability of direct reconstruction techniques to accurately characterise the shape of the central density profile in clusters is very limited, resulting from a decrease in the number of background sources with flexion measurements in the central regions of the cluster. Since direct methods rely on the presence of reliable flexion measurements within each pixel, the ability to measure the central density profile is generally rather limited. Again, as aperture measures involve filtering the signal from a large area, they are less limited by a dearth of sources in the central regions.  Lastly, whilst the direct method needs to be adjusted by a somewhat arbitrary constant shift, the aperture mass statistic is, by construction, independent of this normalisation.  There are often found to be somewhat large differences seen in mass estimates of clusters of galaxies found using strong lensing, weak lensing, or a combination of the two. The flexion aperture mass statistic offers a method to break the degeneracy between these models, as it clearly discriminates between different masses and mass models. As a result of filtering the data, different mass models exhibit significantly different behaviours. Thus, this method offers a unique and novel model-selection method that is complementary to other techniques currently used, and requires no prior assumptions about the underlying mass distribution.  The issue of finding the best-fit mass profile is important to cosmology. Cosmological simulations suggest dark matter haloes should follow a universal (NFW) density profile with concentration parameter anti-correlated with mass. There is much discussion regarding whether observations support these predictions. Thus, the ability to break degeneracies between models is important to cosmology; in this paper, we offer a clear prescription for a method by which this can be achieved.  There is much work remaining to be done, of course. A more finely-sampled parameter space would allow for increased accuracy in model selection, improving the signal to noise in the measurements (by improving the filter functions used, for example) would increase the discriminating power of the $\\fmap$-statistic, and considering the expected radial signal from elliptical lenses or models involving substantial substructure would provide a more realistic sample of model templates from which to select. However, the work presented in this paper represents an excellent first step, and highlights the promise that the aperture mass statistic from flexion shows to be able to discriminate between masses, mass profiles and NFW concentration parameters using only a few combinations of aperture parameters."
    },
    "0910/0910.0135_arXiv.txt": {
        "abstract": "The radio-loud active galactic nucleus in M\\,87 hosts a powerful jet fueled by a super-massive black hole in its center. A bright feature 80\\,pc away from the M\\,87 core has been reported to show superluminal motions, and possibly to be connected with a TeV flare observed around 2005. To complement these studies and to understand the nature of this feature, we analyzed 2\\,cm VLBI data from 15 observing runs between 2000 and 2009. This feature is successfully detected at the milli-Jansky level from 2003 to 2007. Our detections show that its milli-arcsecond structure appears to be extended with a steep spectrum, and no compact or rapidly moving features are observed. Our results do not favor a blazar scenario for this feature. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.0121_arXiv.txt": {
        "abstract": "In this paper, we describe the implementation and performance of GreeM, a massively parallel TreePM code for large-scale cosmological {\\it N}-body simulations.  GreeM uses a recursive multi-section algorithm for domain decomposition. The size of the domains are adjusted so that the total calculation time of the force becomes the same for all processes. The loss of performance due to non-optimal load balancing is around 4\\%, even for more than $10^3$ CPU cores. GreeM runs efficiently on PC clusters and massively-parallel computers, such as a Cray XT4. The measured calculation speed on Cray XT4 is $5\\times 10^4$ particles per second per CPU core, for the case of an opening angle of $\\theta=0.5$, if the number of particles per CPU core is larger than $10^6$. ",
        "introduction": "\\label{sec:Introduction}  The cold dark matter (CDM) model \\citep{White1978, Peacock1999} is widely regarded as the standard theory for the formation and evolution of the universe. According to this model, structure formation in the universe proceeds hierarchically. Small-scale structures form first, and they then merge with each other to form larger-scale structures.  Cosmological {\\it N}-body simulations have been widely used to study nonlinear structure formation in the CDM model.  A number of numerical algorithms for cosmological {\\it N}-body simulations have been proposed.  The $\\rm P^3M$ (Particle Particle Particle Mesh) algorithm, introduced by \\citet{Hockney1981}, is one such algorithm. In the $\\rm P^3M$ algorithm, gravitational interactions between particles are split into short-range and long-range parts. The short-range part is calculated by direct summation (PP part), and the long-range part is calculated by the PM method (PM part), which is accelerated by Fast Fourier Transformation (FFT).  The $\\rm P^3M$ algorithm keeps the advantage of the original PM algorithm, which uses FFT to evaluate the gravitational potential, while improving the spatial resolution by evaluating the short-range force directly. When the universe is close to uniform, the $\\rm P^3M$ method is very fast, since the calculation of FFT is fast and the calculation cost of PP part is small.  However, when the system shows strong clustering, the calculation cost for the PP part becomes very large. Thus, it is not practical to use the $\\rm P^3M$ method for a highly clustered distribution of particles realized in CDM cosmology.  \\citet{Couchman1991} developed the AP3M (adaptive $\\rm P^3M$) algorithm. In this algorithm, the gravitational interactions between particles in high-density regions are split into three or more terms, and only the shortest-range interaction is calculated directly. Intermediate-range terms are evaluated with meshes of different sizes. Conceptually, this AP3M algorithm is simple and efficient. In practice, efficient implementation of the AP3M algorithm on large-scale parallel computers is not easy, and the AP3M algorithm in its fully multi-level form is not widely used. GADGET-2 \\citep{Springel2005} uses two-level hierarchy.  Yet another way to reduce the calculation cost of a high-density region of the $\\rm P^3M$ scheme is to use the tree algorithm instead of direct summation, which is now usually called TreePM \\citep{Xu1995, Bode2000, Bagla2002, Bode2003, Dubinski2004, Springel2005, Yoshikawa2005, Khandai2009}. In this algorithm, the short-range interaction is calculated by the Tree method \\citep{Barnes1986}. The calculation cost per particle of the tree algorithm is almost independent (depends only through the $\\log N$ term) of the local density. Thus, the calculation cost of the PP part of the TreePM algorithm is also nearly independent of the degree of the clustering.  There are two variants of the TreePM algorithm.  One is the TPM algorithm developed by \\citet{Xu1995}.  In this algorithm, the short-range force is calculated using the tree only in high-density regions.  TPM algorithm has been adopted by \\citet{Bode2000} and \\citet{Bode2003}, and is scalable to a large number of CPU processors.  In the other variant \\citep{Bagla2002, Dubinski2004, Springel2005, Yoshikawa2005}, the tree method is applied to the entire simulation box, or to the domain handled by one process. The advantage of this method over the other is the ease of implementation. The tree part of this algorithm (which we call TreePM in this paper, while the other we call TPM) is essentially the same as the treecode for a non-periodic boundary condition, with just two differences. The force law is different, and we need to handle image particles. The PM part is relatively simple anyway. Thus, if one has already developed a parallel treecode for a non-periodic boundary condition, it is straightforward to implement the periodic boundary by adding the calculation code for the PM force. In this case, domain decomposition is based on the need of the tree part, and the PM mesh is used only for the PM force calculation. Since parallelization of the treecode with the non-periodic boundary condition has been very well studied, we can achieve both the ease of implementation, and a good performance, by just using existing algorithms.  On the other hand, in the TPM algorithm, the domain decomposition would usually be based on the PM mesh, since there is no other data structure to rely on. Thus, achieving high efficiency with parallel implementation of TPM requires significant work.  Probably for this reason, many implementations of the parallel TreePM have been reported.  \\citet{Dubinski2004} described GOTPM (Grid-of-Oct-Trees-Particle-Mesh), which is based on one-dimensional slab domain decomposition.  Its performance scales well on hundreds of processes.  \\citet{Springel2005} described GADGET-2 (GAlaxies with Dark matter Gas intEracT), which uses domain decomposition based on a space-filling curve, similar to that used by \\citet{Warren1993}. One advantage of this method is that the tree structure is global, and therefore the force on a particle does not depend on the number of processes used. In some parallel implementations of the tree algorithm, the force on a particle depends on the number of processes used, since the way the force on a particle is calculated is not exactly the same if the number of processes is not the same. This makes the development and validation of the parallel code rather troublesome.  \\citet{Yoshikawa2005} presented the $\\rm P^3M$ and TreePM implementation of cosmological {\\it N}-body code  on GRAPE-5 \\citep{Kawai2000} and GRAPE-6A \\citep{Fukushige2005} systems.  Here, GRAPE \\citep{Sugimoto1990, Makino1998, Makino2003} is used to accelerate the tree part, in the same way as accelerating the tree algorithm for a non-periodic boundary condition \\citep{Makino1991,Makino2004}. We call their code YTPM. It uses one-dimensional (1-D) decomposition, because it is designed for relatively small GRAPE clusters (up to 16 or 32 processors).  If the number of processes is small, 1-D decomposition is okay, simply because other methods do not help in achieving a better parallel performance. However, if the number of processes is large, 1-D decomposition becomes unpractical, because of a increase in the amount of communication and memory to store the particles and tree nodes in boundary layers. Consider the case of a system of $N$ particles simulated on $p$ processes. The number of particles in one process is $N/p$, and the amount of communication (and additional memory requirement) per process per timestep is $O(N^{2/3})$ in the case of 1-D decomposition, and $O[(N/p)^{2/3}]$ in the case of 3-D decomposition. Thus, 3-D decomposition reduces the requirement for communication by a factor of $O(p^{2/3})$, which can be very large, since there are machines with more than $10^5$ processes.  We have developed a parallel TreePM code, which uses domain decomposition based on a recursive multi-section algorithm \\citep{Makino2004}. The recursive multi-section algorithm is similar to the widely used recursive bisection \\citep{Warren1992,Dubinski1996}, but allows 1-D division of an arbitrary number. Thus, it can be used on systems with the number of processes not being powers of two. In addition, we modified the decomposition algorithm so that the load balance becomes practically perfect. Our code, which we call GreeM ({\\it G}RAPE T{\\it  ree}P{\\it M}), is optimized for clusters of GRAPEs or usual PC clusters, which have a relatively poor network performance. Since the communication performance is the limiting factor of the scalability, our code can scale to a very large number of processes, on parallel computers with fast networks, such as Cray XT4.   In section 2, we describe the algorithm used in GreeM.  In section 3, the results of accuracy and performance tests are presented.  Section 4 is for summary and discussion. ",
        "conclusions": "\\subsection{Possible Ways to Improve Accuracy}  As we have shown in subsection 3.1, the total force error of the TreePM method is dominated by the error of the forces from tree nodes at a distance of around $\\eta$. Thus, one possibility of reducing the error is to use distance-dependent opening criterion \\citep{Makino1991}. Since the error in the limit of the uniform density distribution is proportional to $r\\theta^2 g(r/\\eta)$, if we set \\begin{equation} \\theta = \\frac{\\theta_0}{\\sqrt{rg(r/\\eta)}}, \\end{equation} the error is evenly distributed over all nodes, and thus the error should be minimized. It is probably necessary to set some upper-limit value, since with the above criterion alone $\\theta$ would diverge at both ends of $r$.  A more natural approach is to include higher-order terms of expansion. Note that the multipole expansion of force with cutoff is different from that of a pure $1/r$ force, and should include the spacial derivatives of the cutoff function. Thus, its implementation differs from the usual high-order multipole moment. The second-order moment is fairly easy to implement, and can drastically improve the accuracy.  \\subsection{Comparison with Another Code} Here, we compare the performance of GreeM with that of GADGET-2 \\citep{Springel2005}. We discuss only the scalability, because the comparison of the absolute speed does not have much meaning if the hardware used is different. We used the data of figure 19 in \\citet{Springel2005} for a $270^3$ dark matter simulation.  We used a $256^3$ dark matter simulation for GreeM. This comparison is not an exact one since the dark matter distributions and computers were different. However it should give us some useful information.  \\begin{figure}[t] \\centering \\includegraphics[width=10cm]{fig11.eps} \\caption{Speed up of our code and GADGET-2 plotted against the number of CPU cores. The filled circles and open squares show the result of GreeM and that of GADGET-2, respectively [figure 19 in \\citet{Springel2005}]. } \\label{fig:diff} \\end{figure}  Figure \\ref{fig:diff} shows the speed-up factors of our code and GADGET-2 as a function of the number of CPU cores. We can see that the scaling of GreeM is better than that of GADGET-2. The most likely reason for this difference is the difference in the load balance. Even for a very small number of processes, the load imbalance of GADGET-2 is large (as can be seen in their table 1). We achieved a nearly perfect load balance, for an arbitrary number of processes.      \\subsection{Summary}  In this paper, we described our new cosmological {\\it N}-body simulation code, GreeM, which uses the TreePM algorithm, and is optimized for large parallel systems. GreeM achieves a nearly perfect load balance, even for a very large number of cores, resulting in very good scalability.  GreeM runs efficiently on PC clusters, but the scalability is naturally better on parallel computers with high-speed networks. The measured calculation speed on the Cray XT4 is $5 \\times 10^4$ particles per second per CPU core, if the number of particles per CPU core is larger than $5 \\times 10^5$. On a cluster of PCs with quad-core CPU and GbE network, GreeM achieves a similar speed if the number of particles per core is more than $3 \\times 10^6$. Using this code, we have already performed $1600^3$ dark matter simulation on Cray-XT4 \\citep{Ishiyama2009}.  It spent about 0.6 million CPU hours.     \\bigskip We are grateful to Kohji Yoshikawa for providing his parallel TreePM code.  We thank Keigo Nitadori for his technical advices and providing his Phantom GRAPE code.  T.I.  thanks Simon Portegies Zwart and Derek Groen for helpful discussions and advices.  Numerical computations were carried out on a Cray XT4 and a PC cluster at Center for Computational Astrophysics, CfCA, of National Astronomical Observatory of Japan.  T.I. is financially supported by a Research Fellowship of the Japan Society for the Promotion of Science (JSPS) for Young Scientists.  This research is partially supported by the Special Coordination Fund for Promoting Science and Technology (GRAPE-DR project), Ministry of Education, Culture, Sports, Science and Technology, Japan."
    },
    "0910/0910.0251_arXiv.txt": {
        "abstract": "We have conducted a search for ionized gas at 3.6\\,cm, using the Very Large Array, towards 31 Galactic intermediate- and high-mass clumps detected in previous millimeter continuum observations. In the 10 observed fields, 35 HII regions are identified, of which 20 are newly discovered. Many of the HII regions are multiply peaked indicating the presence of a cluster of massive stars. We find that the ionized gas tends to be associated towards the millimeter clumps; of the 31 millimeter clumps observed, 9 of these appear to be physically related to ionized gas, and a further 6 have ionized gas emission within 1'. For clumps with associated ionized gas, the combined mass of the ionizing massive stars is compared to the clump masses to provide an estimate of the instantaneous star formation efficiency. These values range from a few percent to 25\\%, and have an average of 7 $\\pm$ 8\\%. We also find a correlation between the clump mass and the mass of the ionizing massive stars within it, which is consistent with a power law. This result is comparable to the prediction of star formation by competitive accretion that a power law relationship exists between the mass of the most massive star in a cluster and the total mass of the remaining stars. ",
        "introduction": "The study of intermediate ($ 2 $M$_{\\odot} <  $ M$_\\star <8 $M$_{\\odot}$) and high-mass (M$_\\star> 8 $M$_{\\odot}$) star formation has made significant progress in recent years, catching up dramatically with our understanding of low mass star formation \\citep[e.g.,][]{testi03, beuther05, zinnecker07}. In fact, it appears that stars share many similar formation processes over a wide range of stellar masses. As in the case of low-mass star formation, outflows have been observed ubiquitously towards forming intermediate and high-mass stars \\citep[e.g.,][]{fuente01,shepherd960,zhang05}, and there exist examples of stars as massive as early B-type surrounded by disks  \\citep[e.g.,][]{cesaroni07}.  However, a key difference between these regimes is that the reduced Kelvin-Helmholtz contraction timescale of massive protostars results in massive stars reaching the main-sequence while they are still accreting \\citep[][and references therein]{keto06}. Thus, we expect the circumstellar material of these accreting embedded stars to be progressively engulfed by an expanding HII region, and to also be affected by strong stellar winds and radiation pressure. This fact may provide an explanation for why outflows and disks are not commonly detected towards forming O stars \\citep[e.g.,][]{shepherd03, cesaroni07}, as the above processes may be responsible for quickly photoevaporating or dispersing the circumstellar material. However, current observational biases may instead be responsible for the lack of disk detections for the most massive stars.  In this work we aim to provide preliminary studies of a selection of sites containing intermediate- and high-mass star formation, specifically to uncover the presence of ionized gas towards them. We also wish to study the relationship between the star forming gas, traced by millimeter continuum emission from dust, and the ionized gas created by massive stars, traced by radio continuum emission. We will select the most promising objects from this study for follow up with higher resolution observations, to map any outflows or disks towards these sources, and to study how the formation of an HII region affects the material within several hundreds of AU of the star. Therefore, we have selected sources that are within a declination range suitable for future study with Atacama Large Millimeter Array (ALMA) and the Expanded Very Large Array (EVLA).  To fulfill these aims, we have conducted radio continuum observations at 3.6~cm, using the Very Large Array (VLA) of the National Radio Astronomy Observatory \\footnote{The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.}, of 31 clumps detected in previous millimeter continuum observations.   Ten sources were selected from preliminary images from the Bolocam Galactic Plane Survey (BGPS, Aguirre et al. 2009, in preparation), and 5 were selected from \\citet[][hereafter B06]{beltran06}. The remaining 16 sources were observed serendipitously, as their positions lay within the observed VLA 3.6~cm fields.  Note that, in this work, we adopt the terminology that a molecular core produces a single star (or close binary system) while molecular clumps form clusters of stars.  In our study of massive star forming regions, our selected sources are several kiloparsecs away and, thus, we most likely detect clumps forming one or more massive stars along with many lower mass stars.  Section \\ref{sourceselect} presents our selection criteria for the observed millimeter sources. Section \\ref{observation} provides details of our observations and data reduction. Section \\ref{results} presents the observed 3.6~cm continuum images of each field; lists the measured positions, fluxes, and angular sizes of each detected VLA source; compares these results to existing millimeter and mid-IR observations; and presents derived properties for each source. Our discussion is presented in Section \\ref{discussion}. Finally, our conclusions are given in Section \\ref{conclusions}. ",
        "conclusions": "}  Using 3.6\\,cm continuum observations with the VLA, we have surveyed the ionized gas towards 31 intermediate and massive molecular clumps previously observed at millimeter wavelengths. These millimeter clumps were selected from preliminary 1.1\\,mm Bolocam Galactic Plane Survey images, and the 1.2\\,mm observations of \\citet{beltran06}.  In the 10 observed fields, 35 HII regions were identified, 20 of them being newly discovered. The observed HII regions display a wide range of morphologies, and many are multiply peaked, indicating the presence of a cluster of massive stars.  Images comparing the warm dust and ionized gas emission in the 10 observed fields are presented, as well as GLIMPSE survey images of \\textit{Spizter} IRAC emission in these fields. The properties of the millimeter continuum clumps and the HII regions were calculated, and are listed in Tables \\ref{selectsrc} and \\ref{vlaproperties}. There is a large range in the properties of the observed HII regions; their physical sizes extend from $<$0.05\\,pc to 7.88\\,pc, and their spectral types cover B2 to O5. The wealth of information we have provided about the observed millimeter clumps allows them to be followed up in future with more powerful instruments such as the EVLA and ALMA.  By comparing the positions of the millimeter clumps and ionized gas, we have shown that the ionized gas tends to be associated with the millimeter clumps, however this association does not depend on the mass of the clump. Of the 31 millimeter clumps observed, 9 appear to be physically related to ionized gas, and a further 6 have ionized gas emission within 1'.  We also infer ages for several ``types\" of millimeter clump: those with multiply peaked 3.6\\,cm continuum emission, indicating they are clusters, are most likely to pinpoint the most evolved sites of star formation. These are followed by millimeter clumps associated with singly peaked HII regions, which are likely to be a single star or compact cluster, and finally millimeter clumps with no associated ionized gas emission may provide examples of the youngest sources.  Comparing the 1\\,mm Bolocam or SEST images with $^{13}$CO images from the Galactic Ring Survey (GPS), we have shown that the emission from these two tracers is correlated, confirming that distance determinations using $^{13}$CO data can be made for the observed 1\\,mm clumps.  We have compared the masses of massive stars or star clusters to the actual amount of gas available traced by millimeter emission to provide an estimate of the instantaneous star formation efficiency for a clump, giving values ranging from a few percent to 25\\%, with an average of 7 $\\pm$ 8\\%.  We find that the mass of a clump is correlated with the mass of the massive stars which form within it, and that this relationship is consistent with a power law. This is comparable to the results of \\citet{bonnell04}, who find that there is a power-law relationship between the most massive star in a cluster and the combined mass of the remaining stars."
    },
    "0910/0910.1257_arXiv.txt": {
        "abstract": "We report the detection in Ks-band of the secondary eclipse of the hot Jupiter CoRoT--1b, from time series photometry with the ARC 3.5-m telescope at Apache Point Observatory.  The eclipse shows a depth of 0.336$\\pm$0.042 percent and is centered at phase 0.5022$^{+0.0023}_{-0.0027}$, consistent with a zero eccentricity orbit (e $\\cos{\\omega} = 0.0035^{+0.0036}_{-0.0042}$).  We perform the first optical to near--infrared multi--band photometric analysis of an exoplanet's atmosphere and constrain the reflected and thermal emissions by combining our result with the recent 0.6, 0.71, and 2.09 $\\mu$m secondary eclipse detections by \\citet{2009Natur.459..543S}, \\citet{2009arXiv0905.4571G}, and \\citet{2009A&A...501L..23A}. Comparing the multi-wavelength detections to state--of--the--art radiative--convective chemical--equilibrium atmosphere models, we find the near--infrared fluxes difficult to reproduce. The closest blackbody--based and physical models provide the following atmosphere parameters: a temperature $T =2454^{+84}_{-170}$ K, a very low Bond albedo $A_{B} = 0.000^{+0.087}_{-0.000}$, and an energy redistribution parameter $P_n = 0.1$, indicating a small but nonzero amount of heat transfer from the day-- to night--side. The best physical model suggests a thermal inversion layer with an extra optical absorber of opacity $\\kappa_e=0.05$ cm$^2$ g$^{-1}$, placed near the 0.1-bar atmospheric pressure level. This inversion layer is located ten times deeper in the atmosphere than the absorbers used in models to fit mid-infrared Spitzer detections of other irradiated hot Jupiters. ",
        "introduction": "\\label{sec:intro}  Space--based detections of hot Jupiter atmospheres have flourished in recent years.  The Spitzer Space Telescope has successfully detected thermal emission from several planets at wavelengths longer than 3.6 $\\mu$m (e.g.~\\citealt{2005ApJ...626..523C,2005Natur.434..740D, 2007MNRAS.378..148D,2007Natur.447..183K,2008ApJ...673..526K, 2007Natur.447..691H,2008ApJ...684.1427M,2009arXiv0906.1293M}), while near--infrared (1.5--2.5 $\\mu$m) observations with the Hubble Space Telescope have found evidence for water, carbon monoxide, and carbon dioxide in the dayside spectrum of the exoplanet HD 189733b \\citep{2009ApJ...690L.114S}. These detections have been made during secondary eclipses (when the planets pass behind their host stars). Important atmospheric absorption signatures have been detected by transit observations with these space telescopes as well: sodium in HD 209458b \\citep{2002ApJ...568..377C} and methane in HD 189733b \\citep{2008Natur.452..329S} with Hubble, and water vapor in HD 189733b with Spitzer \\citep{2007Natur.448..169T}.  The CoRoT mission recently joined these space--borne successes by detecting phase brightness variations and combined thermal and reflected emission during secondary eclipses of the exoplanets \\corot \\citep{2009Natur.459..543S,2009A&A...501L..23A} and CoRoT--2b \\citep{2009arXiv0906.2814A} in an optical broadband window centered at about 0.6 $\\mu$m (see Table~\\ref{table1}).  At the same time that these detections have provided very valuable insights into the atmospheric physics of irradiated hot Jupiters, they have also revealed some perplexing findings. Some of the planets show a strong temperature contrast between their day and night sides (e.g.~\\citealt{2006Sci...314..623H,2009Natur.459..543S}), while others appear to have a more efficient redistribution of incident energy (e.g.~\\citealt{2007Natur.447..183K}). Another striking discovery is an apparent bifurcation of hot Jupiters into two classes based on the presence or absence of a thermal inversion layer (e.g.~\\citealt{2003ApJ...594.1011H,2008ApJ...682.1277B,2008ApJ...678.1419F}). Along with a wider--than--expected range of exoplanet radii (see e.g.~\\citealt{2008A&A...482L..17B}), these are currently the key unsolved questions in the study of irradiated hot Jupiter atmospheres.  Resolving these questions will require observations from the optical to the infrared wavelength regimes. Ground-based observations have just recently started to reach the sensitivity necessary to directly detect hot Jupiters. The first two detections were announced in early 2009 at optical and near--infrared wavelengths:  \\ogle in z'-band \\citep{2009A&A...493L..31S}, and TrES-3 in K-band \\citep{2009A&A...493L..35D}. The third ground--based result was the detection of an eclipse of \\corot at 2.09 $\\mu$m \\citep{2009arXiv0905.4571G}. This last detection makes \\corot the first exoplanet to have thermal emission measured at both optical and near--infrared wavelengths.  Here we report the fourth ground-based detection of thermal emission from an exoplanet, and the fourth detection of CoRoT--1b, this time in Ks-band (2.15 $\\mu$m). We combine our results with the other three recent detections of \\corot by \\citet{2009Natur.459..543S}, \\citet{2009arXiv0905.4571G}, and \\citet{2009A&A...501L..23A} to produce the first multi--color analysis of the atmospheric spectrum of an exoplanet between 0.5 and 2.2 $\\mu$m.  Section \\ref{sec:obs} describes the observations.  In section \\ref{sec:raa} we detail our reduction and analysis steps to detect the signal from the planet. Section \\ref{sec:models} compares the detected signals to the predictions made by state-of-the-art planetary atmosphere models. The results are then discussed and summarized in sections \\ref{sec:discussion} and \\ref{sec:summary}. ",
        "conclusions": "\\label{sec:discussion}  The detections of exoplanets via transits and secondary eclipses have produced many of the most valuable insights to their physical properties, but also yielded some puzzling results.  {Among these is the apparent division of the population of hot Jupiters into two distinct classes based on their atmospheric properties: one group that appears cooler, with water and methane absorption bands and more efficient energy redistribution, and another with higher levels of thermal emission, strong day--night contrasts, and lacking the expected absorption bands. The dominant explanation for this dichotomy is the presence in many planets of a thermal inversion layer, i.e.~a hot stratosphere caused by extra absorbers of optical light \\citep{2003ApJ...594.1011H,2008ApJ...682.1277B,2008ApJ...678.1419F}.  The exact nature of the absorbers is a difficult question still being investigated. \\citet{2009arXiv0902.3995S} recently argued that vanadium oxide is not likely to fulfill this role, and that the previously favored titanium oxide would require unusually high levels of macroscopic mixing to remain in the upper atmosphere.  S$_2$, S$_3$, and HS compounds, as absorbers of optical and ultraviolet light, have also recently been both considered and questioned as causes of the thermal inversion \\citep{2009AAS...21430601Z,2009arXiv0903.1663Z}. Consistent modeling of the infrared secondary eclipse spectra of six planets by \\citet{2008ApJ...678.1436B} suggest that the presence of the necessary absorber may depend not only on the incident stellar radiation, but also on planetary metallicity and surface gravity.  From the combination of the {\\bf four} optical to near--infrared detections we have detailed in this paper, it appears that \\corot clearly falls into the class of hot Jupiters with a thermal inversion. The best--fit physical model to the combined dataset requires a small redistribution parameter P$_n$ = 0.1 and an extra optical absorber with flat opacity $\\kappa$ = 0.05 cm$^2$/g. While we cannot determine the identity of this absorber, the model constrains its altitude, requiring absorption a factor of 10--100 times deeper in the atmosphere than suggested by previous results for other exoplanets. This new constraint also highlights the power of combined optical and near-infrared photometry.  \\corot reflects some other surprising trends as well. Hot Jupiters tend to have very low albedos \\citep{2000ApJ...538..885S,2006ApJ...646.1241R,2008ApJ...682.1277B}, and our findings place this planet's atmosphere in agreement with that trend. From the best--fit physical model derived in Section \\ref{sec:models_TheoreticalModels}, we obtained an estimate of the planet's geometric albedo at the wavelengths probed by the CoRoT detections (0.4 to 1.0$\\mu$m) to be A$_g$ = 0.05$\\pm$0.01. Assuming a wavelength-independent Lambert sphere (see \\citealt{2007ApJ...667L.191L}), this corresponds to a Bond albedo A$_B$=0.075$\\pm$0.015.  \\corot is also notable for its extremely large radius (1.45 R$_{Jup}$, \\citealt{2009arXiv0905.4571G}); evolutionary models for a hot Jupiter of its mass, age, and irradiation predict a radius of only 0.94 - 1.18 R$_{Jup}$ \\citep{2007ApJ...659.1661F}. This places it towards the upper end of the wide distribution of radii that has been seen among the population of transiting exoplanets.  While small--radius planets can be modeled with a larger, denser core, the inflated sizes of planets like \\corot provide quite a challenge to planetary models. Several explanations have been proposed, including heat retained by enhanced atmospheric opacities \\citep{2007ApJ...668L.171B}, deposition of kinetic wind energy in the upper atmosphere \\citep{2008ApJ...682..559S}, and significant tidal heating caused by orbital eccentricity or a rapidly-rotating star \\citep{2003ApJ...592..555B,2009ApJ...698L..42G,2009AAS...21430602M,2009arXiv0902.3998I}. Though \\citet{2009arXiv0905.4571G} measured an eccentricity $e$=0.071$^{+0.042}_{-0.028}$ for CoRoT--1b's orbit, the eccentricity we measure ($e\\cos{\\omega}$=0.0035$^{+0.0036}_{-0.0042}$) based on its mid--eclipse phase is consistent within its errors with a circular orbit. Thus, there is no evidence that tidal heating derived from a currently eccentric orbit contributes to the energy budget of CoRoT--1b, although intense heating in the recent past cannot be excluded (e.g.~\\citet{2009ApJ...702.1413M}). In each of the diverse models, the radius of the planet is determined by the pressure--temperature structure of the atmosphere and the balance between the energy input and output. In order to understand and improve the models, we must first constrain the energy balance by understanding the basic atmospheric properties, such as chemical abundances, pressure--temperature profiles, and energy redistribution.  With our detection in Ks joining the 0.6, 0.71 and 2.09 $\\mu$m measurements, we are beginning to get a more detailed picture of the atmospheric properties of CoRoT--1b. However, there are still sizeable gaps in the observed spectrum, and even the most advanced current models have difficulty explaining the high flux levels in the 2-$\\mu$m window. It is crucial to add further detections in both narrow and broad bands, in the optical and near-infrared, and a larger target sample to provide stronger constraints to the planetary atmosphere models.      -- We directly detect thermal emission from CoRoT--1b in Ks-band, determining the planet-to-star flux ratio to be 0.336$\\pm$0.042\\%.  -- Using simple blackbody--based models, we find the best fit to be a blackbody of 2454$^{+84}_{-170}$~K, and confidently rule out both a high albedo and efficient day--night heat redistribution.  -- Using realistic atmosphere models, we find the need for a thermal inversion layer and an extra absorber near the 0.1--bar level, deeper in the atmosphere than the Spitzer mid--infrared data suggested for other hot jupiters.  -- Both the blackbody models and physical atmospheric models agree on a small but non-zero amount of heat redistribution, and a Bond albedo less than 0.09.  In short, the combined optical to near--infrared photometry of \\corot has allowed us to independently constrain the reflected light and thermal emission, and has revealed a very hot, low--albedo planet with a large day--night contrast and a prominent temperature inversion."
    },
    "0910/0910.4969_arXiv.txt": {
        "abstract": "We simulate the possible emission from a disk perturbed by a recoiling super-massive black hole.  To this end, we study radiation transfer from the system incorporating bremsstrahlung emission from a Maxwellian plasma and absorption given by Kramer's opacity law modified to incorporate blackbody effects.  We employ this model in the radiation transfer integration to compute the luminosity at several frequencies, and compare with previous bremsstrahlung luminosity estimations from a transparent limit (in which the emissivity is integrated over the computational domain and over all frequencies) and with a simple thermal emission model.  We find close agreement between the radiation transfer results and the estimated bremsstrahlung luminosity from previous work for electromagnetic signals above $10^{14}$~Hz.  For lower frequencies, we find a self-eclipsing behavior in the disk, resulting in a strong intensity variability connected to the orbital period of the disk. ",
        "introduction": "The near possibility of detecting gravitational waves to study astrophysical systems has spurred efforts to understand such systems theoretically to aid in the detection and interpretation of obtained signals (see, for instance, \\cite{Aylott:2009tn,Aylott:2009ya,Shoemaker:2008pe,Baumgarte:2006en, Pan:2007nw,Ajith:2007kx} and references cited therein for further details and examples). Systems that radiate strongly in both electromagnetic and gravitational wave bands are particularly interesting since the combined information will provide unprecedented access to a number of rich phenomena (e.g.~\\cite{Haiman:2008zy,Bloom:2009vx,Blecha:2008mg}.) Electromagnetic radiation from compact objects can be (partially or completely) absorbed and scattered by dust and other surrounding material.  Gravitational radiation, on the other hand, is generated in the central engine, and these waves do not interact with dust or material in the surrounding environment.   One such system involves the collision of a binary black hole within a circumbinary disk \\cite{2002ApJ...567L...9A,2003MNRAS.340..411L,Milosavljevic:2004cg}. For a generic black hole binary, the gravitational waves are radiated asymmetrically.  As these waves carry energy and momentum, the resultant black hole can have a kick or recoil velocity~\\cite{kick1,kick3,kick5,kick7,kick9,kick10}, which in turn will induce shocks in the disk. A second source of perturbations in the disk comes from the loss of gravitational mass with the emission of gravitational radiation, weakening the gravitational potential felt by the disk.  These effects, as discussed in~\\cite{kickdisk3,kickdisk1,kickdisk2}, can induce possible emissions and some of these have been investigated in part in~\\cite{kickdisk1,kickdisk2,massloss,kickdisk,Corrales:2009nv}.  These works study the dynamics of the disks when either mass reduction (due to the energy lost by the system) or recoil velocity (due to the asymmetric flux of gravitational radiation) effects are included and present preliminary estimates of  emissions from the system with the goal of understanding possible observations in the electromagnetic spectrum.  In the present work, we further investigate electromagnetic radiation from shocked disks by directly computing the electromagnetic emission using radiation transfer.  We build on our earlier work on disks~\\cite{kickdisk}, hereafter referred to as Paper~I, where we followed the evolution of disks when the central black hole was given different recoil velocities. We use this data for the evolution of  disks to calculate the electromagnetic spectrum using radiation transfer. In this method, we assume simple models for the emission and absorption of electromagnetic radiation and solve for the intensities by integrating the radiation transfer equation along geodesics of the spacetime. We examine multiple radiation models to investigate how different choices in the emission models affect the time-dependent luminosity. One model assumes only classical thermal emission, and a second model incorporates both bremsstrahlung and thermal emission according to optical depth. Finally, we also compare our results with the simple bremsstrahlung estimate that we obtained in Paper~I. ",
        "conclusions": ""
    },
    "0910/0910.3855_arXiv.txt": {
        "abstract": "In this paper we place observational constraints on the interaction and spatial curvature in the holographic dark energy model. We consider three kinds of phenomenological interactions between holographic dark energy and matter, i.e., the interaction term $Q$ is proportional to the energy densities of dark energy ($\\rho_{\\Lambda}$), matter ($\\rho_{m}$), and matter plus dark energy ($\\rho_m+\\rho_{\\Lambda}$). For probing the interaction and spatial curvature in the holographic dark energy model, we use the latest observational data including the type Ia supernovae (SNIa) Constitution data, the shift parameter of the cosmic microwave background (CMB) given by the five-year Wilkinson Microwave Anisotropy Probe (WMAP5) observations, and the baryon acoustic oscillation (BAO) measurement from the Sloan Digital Sky Survey (SDSS). Our results show that the interaction and spatial curvature in the holographic dark energy model are both rather small. Besides, it is interesting to find that there exists significant degeneracy between the phenomenological interaction and the spatial curvature in the holographic dark energy model. ",
        "introduction": "Observations of type Ia supernovae (SNIa) \\cite{Riess}, cosmic microwave background (CMB) \\cite{spergel} and large scale structure (LSS) \\cite{Tegmark} all indicate the existence of mysterious dark energy driving the current accelerated expansion of universe, and lots of efforts have been made to understand it. The most important theoretical candidate of dark energy is the Einstein's cosmological constant $\\lambda$, which can fit the observations well so far, but is plagued with the famous fine-tuning and cosmic coincidence problems \\cite{Weinberg}. Many other dynamical dark energy models have also been proposed in the literature, such as quintessence \\cite{quint}, phantom \\cite{phantom}, $k$-essence \\cite{k}, tachyon \\cite{tachyonic}, hessence \\cite{hessence}, Chaplygin gas \\cite{Chaplygin}, and Yang-Mills condensate \\cite{YMC}, etc.  Actually, the dark energy problem may be in essence an issue of quantum gravity \\cite{Witten:2000zk}. However, by far, we have no a complete theory of quantum gravity, so it seems that we have to consider the effects of gravity in some effective quantum field theory in which some fundamental principles of quantum gravity should be taken into account. It is commonly believed that the holographic principle \\cite{Holography} is just a fundamental principle of quantum gravity. Based on the effective quantum field theory, Cohen et al. \\cite{cohen99} pointed out that the quantum zero-point energy of a system with size $L$ should not exceed the mass of a black hole with the same size, i.e. $L^3\\Lambda^4\\leq LM_{Pl}^2$, where $\\Lambda$ is the ultraviolet (UV) cutoff of the effective quantum field theory, which is closely related to the quantum zero-point energy density, and $M_{Pl}\\equiv 1/\\sqrt{8\\pi G}$ is the reduced Planck mass. This observation relates the UV cutoff of a system to its infrared (IR) cutoff. When we take the whole universe into account, the vacuum energy related to this holographic principle can be viewed as dark energy (its energy density is denoted as $\\rho_{\\Lambda}$ hereafter). The largest IR cutoff $L$ is chosen by saturating the inequality, so that we get the holographic dark energy density \\begin{equation} \\label{eq:rhode0}\\rho_{\\Lambda}=3c^2M_{Pl}^2L^{-2} \\end{equation} where $c$ is a numerical constant characterizing all of the uncertainties of the theory, and its value can only be determined by observations. If we take $L$ as the size of the current universe, say, the Hubble radius $H^{-1}$, then the dark energy density will be close to the observational result. However, Hsu \\cite{hsu04} pointed out this yields a wrong equation of state for dark energy. Subsequently, Li \\cite{li04} suggested to choose the future event horizon of the universe as the IR cutoff of this theory. This choice not only gives a reasonable value for dark energy, but also leads to an accelerated universe. Moreover, the cosmic coincidence problem can also be explained successfully in this model. Most recently, a calculation of the Casimir energy of the photon field in a de Sitter space is performed \\cite{lmp}, and it is a surprising result that the Casimir energy is indeed proportional to the size of the horizon (the usual Casimir energy in a cavity is inversely proportional to the size of the cavity), in agreement with the holographic dark energy model.  The holographic dark energy model has been studied widely \\cite{holode1,holode2,holodeobs,holoobs09,intholo,holoissue,healworld}. By far, various observational constraints on this model all indicate that the parameter $c$ is less than 1, implying that the holographic dark energy would lead to a phantom universe with big rip as its ultimate fate \\cite{holodeobs,holoobs09}. Nevertheless, the appearance of the cosmic doomsday will definitely ruin the theoretical foundation of the holographic dark energy model, namely, the effective quantum field theory; see Ref.~\\cite{healworld} for detailed discussions. To evade this difficulty, one may consider the extra dimension effects in this model \\cite{healworld}: The high energy correction from the extra dimensions could erase the big-rip singularity and leads to a de Sitter finale for the holographic cosmos. Another way of avoiding the big rip is to consider some phenomenological interaction between holographic dark energy and matter \\cite{intholo,holoissue}. With the help of the interaction, the big rip might be avoided due to the occurrence of an attractor solution in which the effective equations of state of dark energy and matter become identical in the far future. However, it should be pointed out that in order to avoid the phantom-like universe the phenomenological interaction term should be tuned to satisfy some specific condition, i.e., the interaction must be strong enough; for detailed discussions see Ref.~\\cite{holoissue}. The only way of acquiring the knowledge of the interaction strength is from the observational data fitting. In this paper, we will explore the possible phenomenological interactions between holographic dark energy and matter by using the latest observational data.  Another important mission of this paper is to probe the spatial curvature of the universe in the holographic dark energy model. Actually, the current observational data are not accurate enough to distinguish between the dynamical effects of dark energy and spatial curvature of the universe, owing to the degeneracy between them. Indeed, in fitting the equation of state of dark energy, $w(z)$, the inclusion of $\\Omega_{k0}$ as an additional parameter would dilute constraints on $w(z)$. While most inflation models in which the inflationary periods last for much longer than 60 $e$-folds predict a spatially flat universe, $\\Omega_{k0}\\sim 10^{-5}$, the current constraint on $\\Omega_{k0}$ is three orders of magnitude larger than this inflation prediction. As argued in Ref.~\\cite{Clarkson:2007bc}, the studies of dark energy, and in particular, of observational data, should include $\\Omega_{k0}$ as a parameter to be fitted alongside the $w(z)$ parameters. So, in this paper, we will constrain $\\Omega_{k0}$ in the holographic dark energy model in light of the latest observational data. Of course, we will also let the interaction and spatial curvature simultaneously be free parameters, and explore the possible degeneracy between them. ",
        "conclusions": "In this paper we consider a sophisticated holographic dark energy model where the interaction and spatial curvature are both involved. The consideration of interaction between dark energy and matter in the holographic dark energy is rather popular, since the interaction not only can help to alleviate the cosmic coincidence problem but also can be used to avoid the future big-rip singularity caused by $c<1$. In addition, the consideration of spatial geometry in the holographic dark energy is also necessary, because usually there exists some degeneracy between the spatial curvature and the dynamics of dark energy. The aim of this paper is to probe the possible interaction and spatial curvature in the holographic dark energy model in light of the latest observational data.  We considered three kinds of phenomenological interactions between holographic dark energy and matter, i.e., the interaction term $Q$ is proportional to the energy densities of dark energy ($\\rho_{\\Lambda}$), matter ($\\rho_m$) and dark energy plus matter ($\\rho_\\Lambda+\\rho_m$). For the observational data, we used the SNIa Constitution data, the shift parameter $R$ of the CMB from the WMAP5, and the BAO parameter $A$ from the SDSS. We separately considered the following cases: opening interaction but closing spatial curvature; opening spatial curvature but closing interaction; simultaneously opening interaction and spatial curvature. Our results show that when separately considering the interaction and spatial curvature, the interaction and spatial curvature in the holographic dark energy model are both rather small. In other words, the observations favor a non-interacting holographic dark energy with flat geometry. When considering both the interaction and the spatial curvature in the holographic dark energy model, it is interesting to find that there exists remarkable degeneracy between them. On the whole, according to the current observational data, the holographic dark energy model without interaction and spatial curvature is still more favored."
    },
    "0910/0910.4191_arXiv.txt": {
        "abstract": "We present optical/near-IR integral field unit (IFU) observations of a gas pillar in the Galactic H\\two\\ region NGC 6357 containing the young open star cluster Pismis 24. These observations have allowed us to examined in detail the gas conditions of the strong wind-clump interactions taking place on its surface.  By accurately decomposing the H$\\alpha$ line profile, we identify the presence of a narrow ($\\sim$20~\\kms) and broad (50--150~\\kms) component, both of which we can associate with the pillar and its surroundings. Furthermore, the broadest broad component widths are found in a region that follows the shape of the eastern pillar edge. These connections have allowed us to firmly associate the broad component with emission from ionized gas within turbulent mixing layers on the pillar's surface set up by the shear flows of the winds from the O stars in the cluster. We discuss the implications of our findings in terms of the broad emission line component that is increasingly found in extragalactic starburst environments. Although the broad line widths found here are narrower, we conclude that the mechanisms producing both must be the same. The difference in line widths may result from the lower total mechanical wind energy produced by the O stars in Pismis 24 compared to that from a typical young massive star cluster found in a starburst galaxy.  The pillar's edge is also clearly defined by dense ($\\lesssim$5000~\\cmt), hot ($\\gtrsim$20\\,000~K), and excited (via the [N\\two]/H$\\alpha$ and [S\\two]/H$\\alpha$ ratios) gas conditions, implying the presence of a D-type ionization front propagating into the pillar surface. Although there must be both photoevaporation outflows produced by the ionization front, and mass-loss through mechanical ablation, we see no evidence for any significant bulk gas motions on or around the pillar. We postulate that the evaporated/ablated gas must be rapidly heated before being entrained. ",
        "introduction": "\\label{sect:intro}  \\subsection{Broad emission features} Direct evidence of interactions between stellar winds and their surroundings is relatively sparse. Observations of H$_{2}$ and CO absorption lines on sight-lines towards massive stars in our Galaxy show comparatively broad underlying components \\citep*{falgarone90, hartquist92}. These were interpreted as arising in turbulent surface layers on cold dense molecular clouds surrounding the stars resulting from the impact of the stellar winds. Broad underlying components have also been observed in the optical emission lines from more energetic environments in nearby giant H\\two\\ regions \\citep[e.g.][]{roy92, chuken94, yang96} and starburst galaxies (e.g.~\\citealt{homeier99, marlowe95, mendez97}; \\citealt*{sidoli06}; \\citealt{vanzi06}). However, due to mismatches in spectral and spatial resolution and in the specific environments observed, the nature of the energy source for these broad optical lines has been contested. Recent detailed integral field unit (IFU) studies of the ionized interstellar medium (ISM) in the nearby starburst galaxies NGC 1569 and M82, however, have shed a considerable amount of light on this problem \\citep{westm07a, westm07b, westm08, westm09a, westm09b}.  By mapping out the properties of the individual emission line components in NGC 1569 (including a broad underlying component with FWHM $\\sim$ 100--350~\\kms), \\citet{westm07a} identified a number of correlations that allowed them to determine the likely origin of the broad component. They concluded that, like in the CO line case, the evaporation and/or ablation of material from cool interstellar gas clouds caused by the impact of the high-energy photons and fast-flowing cluster winds \\citep{pittard05} produces turbulent mixing layers (TMLs) on the surfaces of the clouds \\citep{begelman90, slavin93, binette99, binette09} from which the emission arises. \\citet{westm07c, westm09a, westm09b} showed that this explanation is also applicable to the broad emission lines detected in M82. They argued that since M82's high pressure ISM is highly fragmented with many small clouds well mixed in with the star clusters, there are many cloud surfaces with which the copious ionizing photons and fast winds can interact and this can explain the prominent broad line emission.  \\citet{binette09} modelled the H$\\alpha$ and [O\\three]$\\lambda$5007 broad emission components observed in the extragalactic giant H\\two\\ region NGC 2363 with a TML model. By using an approximate treatment of turbulence \\citep[the mixing length approach of][]{canto91}, they avoided dealing with the details of the turbulent energy cascade, while still being able to predict the line kinematics. The \\textsc{Mappings Ic} code \\citep{ferruit97} was used to compute the radiative transfer and thus the line emissivities. In this model the line width is created by the velocity gradient in the mixing layer. A limitation of this approach is that the turbulent velocity that should be generated in the TML is not included and neither are the non-equilibrium ionization (NEI) effects that are expected when cooler gas mixes rapidly with hotter gas. This model is therefore only a good approximation if the mixing occurs fairly slowly and the ionization is dominated by photoionization. The model of \\citet{slavin93} includes NEI effects, but does not predict line profiles since the dynamics in the layer are not modelled.  The only way to connect the low-energy molecular line cases to the higher-energy optical emission line cases is to model the physical process using the correct initial conditions. We have begun to do this using a hydrodynamical code that includes NEI effects, but we must first confirm that the physical mechanism producing the broad component in both cases is wind-driven TMLs, and we must have reliable measurements of the gas conditions within these layers. Thus we have undertaken a study of a nearby Galactic analogue to the type of environments found in optical studies of extragalactic giant H\\two\\ regions and starburst galaxies. Although the energies and ISM conditions found in starburst galaxies are not easily found within our Galaxy, the advantage of proximity allows discrete superimposed kinematical components to be clearly resolved and unambiguously disentangled from other possible line broadening mechanisms.  \\begin{figure*} \\centering \\includegraphics[width=\\textwidth]{finder_astroph.eps} \\caption{Composite image of NGC 6357: the upper half comprises an ACS+WFPC2 F658N (H$\\alpha$) composite. Since this image does not cover the stars of the cluster, we have extended the figure southwards by including part of the \\textit{HST}/ACS WFC F850W ($\\sim$$I$-band) image. The ARGUS IFU position is outlined with a solid rectangle ($11.5\\times 7.3$ arcsec), the two main ionizing stars \\citep{walborn02, bohigas04} and the ionization front are labelled, and the approximate projected distance from the stars to the tip of the pillar is indicated. \\textit{Inset left, top:} zoom-in showing the outline of the IFU position on the pillar. The image intensity scale has been stretched to enhance the protrusion indicated by the arrow and discussed in the text. \\textit{Inset left, bottom:} integrated H$\\alpha$ flux (C1+C2+C3) map as seen through the IFU. H$\\alpha$ line profiles and Gaussian fits from the two spaxels indicated are shown in the \\textit{top right inset}, where the flux units are in an arbitrary but relative scale.} \\label{fig:finder} \\end{figure*}  \\subsection{NGC 6357}  The analogue we have chosen is a gas pillar that forms part of the youngest H\\two\\ region (G353.2+0.9) in NGC 6357 (RCW 131, W 22, Sh-2 11), a large H\\two\\ region/molecular cloud complex in the Sagittarius spiral arm of the Milky Way. G353.2+0.9 is ionized by the open cluster Pismis 24. This cluster has a stellar population a few times that of the Orion star cluster and a central stellar density similar to that of the Galactic young massive cluster NGC 3603 \\citep{wang07}. Pismis 24 is still too young to have had any supernova explosions \\citep{wang07}; this is in-keeping with its O-star population, which includes an O3.5 If (Pismis 24-1) and O3.5 III(f*) (Pismis 24-17) star \\citep{walborn02}, each of $>$100~\\Msol\\ \\citep{maiz07}.  The conclusion of a number of individual radio and IR studies of NGC 6357 is that most of the molecular material is located either behind the H\\two\\ region and/or north of the sharp ionization front clearly visible in Fig.~\\ref{fig:finder} \\citep{massi97}. The expansion of the H\\two\\ region has caused gas to accumulate in a (large) interim zone into which the ionization front is proceeding; this has triggered the formation of several more (less massive) stars within the cloud \\citep{bohigas04}. The molecular gas observations suggest that NGC 6357 represents a late-stage blister-type H\\two\\ complex viewed face-on, meaning that we are seeing the top ionized face of a much larger molecular cloud \\citep{massi97}, and that the pillar we have focussed on must project up from this structure towards the observer. Since the most highly excited regions of the nebula are facing the cluster, it is thought that the primary ionizing sources must be the two O3.5 stars in the Pismis 24 cluster. This is confirmed by ionization calculations that show that the exciting stars must be of spectral type earlier than O5 \\citep{bohigas04}. The location of these stars and their projected distance to the gas pillar and ionization front (3.5~pc) are indicated in Fig.~\\ref{fig:finder}.  The mechanisms that lead to the formation of ``elephant trunks'' in H\\two\\ regions is still a matter of debate \\citep[e.g.][]{mizuta06, gritschneder09}. However, the fact that these pillar-shaped nebulae usually point towards and exhibit velocity gradients in the direction of a nearby OB star, and show signs of compression and/or star formation at their tips, suggests a strong interaction with the winds and radiation from these sources \\citep{pound98, sugitani02, gahm06, mizuta06, gritschneder09}. Our targeted pillar in NGC 6357 both points directly towards (in projection) the aforementioned O stars in Pismis 24 (Fig.~\\ref{fig:finder}), and has been detected in both CO and HCO$^{+}$ implying that it contains very high density gas \\citep[$>$$10^{5}$--$10^6$~\\cmt;][]{massi97}. The pillar is also known to house two compact radio sources thought to be young stellar objects, including one at its apex \\citep[IRS 4;][]{persi86, felli90, wang07}.   In this paper we present optical/near-IR VLT/FLAMES-ARGUS IFU observations of the tip of this gas pillar (Section~\\ref{sect:obs}). In these observations we have identified multiple broad and narrow emission line components. This has allowed us to examine in detail the kinematics (radial velocity and FWHM) of each H$\\alpha$ line component, and the spatially resolved nebular properties (densities, temperatures, excitations) of this region (Section~\\ref{sect:maps}), and use these insights to discuss the wind and ionization interactions taking place on the pillar's surface (Section~\\ref{sect:disc}). A summary of these findings and discussion is contained in Section~\\ref{sect:conc}. We adopt a distance to the NGC 6357 complex of 2.56~kpc \\citep{massey01}, meaning 10 arcsec $\\sim$ 0.62~pc. ",
        "conclusions": "\\label{sect:conc} In this study we have examined in detail the gas conditions in the surface interaction layers of a gas pillar in NGC 6357 ionized by the star cluster Pismis 24 using high-resolution optical/near-IR IFU observations. Our observations show that the gas on the pillar's edge is denser, hotter, and more highly excited than its surroundings, providing clear evidence that the incoming radiation and winds from the cluster stars are strongly affecting the state of the ionized gas in these surface layers.  That we have been able to accurately decompose the H$\\alpha$ line into multiple components has allowed us to examine the gas dynamics of this region in great detail. One of the primary aims of this study was to look for the presence of a broad emission component, and to investigate whether this component originates in turbulent mixing layers \\citep[TMLs][]{slavin93, binette09} on the surface of the pillar set up by the shear flows of the winds from the young massive stars of the cluster Pismis 24, since this is the scenario suggesed by \\citet[][see also \\citealt{westm07b, westm07c, westm09a, westm09b}]{westm07a} to explain the broad emission components found in the starburst galaxies NGC 1569 and M82. On the pillar we find clear evidence of both a narrow ($\\sim$20~\\kms) and broad (50--150~\\kms) component; away from the pillar we also see a much brighter, redshifted narrow component. The similarity of the broad C2 component velocities to the velocity of the narrow line component emitted at the location of the pillar leads us to conclude that the turbulent, C2-emitting gas must be associated with the pillar structure, whereas the redshifted narrow component originates in the background gas. That the velocities of the components associated with the pillar are consistent with those of the CO gas, and the fact that the broadest C2 line widths are found to follow the shape of the eastern pillar edge, lends further support to the connection between C2 and turbulent ionized gas within the pillar surface layers. We can therefore assert that the mechanical energy of the stellar winds impacting on the pillar surface is driving a turbulent mixing layer, and it is from this layer that the broad component is being emitted.  The pillar's edge is also clearly defined by dense ($\\lesssim$5000~\\cmt), hot ($\\gtrsim$20\\,000~K), and excited (via the [N\\two]/H$\\alpha$ and [S\\two]/H$\\alpha$ ratios) gas conditions, implying the presence of a D-type ionization front propagating into the pillar surface. Although there must be both photoevaporation outflows produced by the ionization front, and mass-loss through mechanical ablation (stripping), we see no evidence for any significant bulk gas motions on or around the pillar. We postulate that the evaporated/ablated gas must be rapidly heated before being entrained (i.e.\\ the gas heating time is fast compared to outflow time).  Although the location of the broadest H$\\alpha$ C2 emission follows the shape of the eastern pillar edge, it continues from the southern tip along a diagonal strip extending in the south-west direction. This extension is coincident with a faint protrusion to the tip of the pillar seen in the \\textit{HST}/ACS F658N image. We conclude that it is either a part of the original molecular cloud that remains to be eroded or that it has been formed through a hydrodynamical instability in the outflowing material, and the geometry is such that it is shielding the western edge of the main pillar body from the oncoming winds therefore preventing the development of a TML on its surface.  Finally we have discussed the implications of our findings to what we have previously observed in starburst galaxies. The C2/C1 flux ratio is, in general, similar to what we have found in these energetic extragalactic environments, suggesting that the TMLs in both cases are of similar importance. That the C2 widths are narrower (50--150~\\kms\\ \\textit{vs.} 150--350~\\kms), however, could result from the fact that the Pismis 24 cluster contains only a few O stars and is too young to have had any supernovae, so the total mass flux and therefore mechanical wind energy is at least two orders of magnitude lower than that from a typical young massive star cluster such as those found in starburst galaxies. This effect results in a lower level of turbulence in the surface layers of the pillar.  To further investigate the relationships between the different gas components within the pillar, its surface and the surrounding gas more fully, it would be important to derive excitation, density and temperature diagnostics for each of the individual line components (C1, C2 and C3). This we have not been able to do since our observations of the faintest lines required for these diagnostics were not of a sufficient S/N. Thus, we intend to obtain deeper observations of this object in the future to build on the work presented here."
    },
    "0910/0910.4677.txt": {
        "abstract": "Common wisdom associates all the unraveled and theoretically challenging aspects of gravity with its UV-completion. However, there appear to be few difficulties afflicting the effective framework for gravity already at low energy, that  are likely to be detached from the high-energy structure. Those include the black hole information paradox, the cosmological constant problem and the rather involved and fine tuned model building required to explain our cosmological observations.  I review some on-going research that aims to generalize and extend the low-energy framework for gravity. In a quantum informational fashion, regions of space at a given time are treated and described as quantum subsystems, rather than submanifolds. The idea is to define a region of space through the quantum degrees of freedom of the matter fields ``living therein\". I show how the correspondence sub-system/sub-manifold is realized in standard semi-classical gravity (``standard localization\") and discuss the implications of alternative localization schemes. By exploiting further the subsystem description, I then consider the possibility that the usual GR metric manifold description might break down in the infra-red. This has implications on the description of the Universe on the largest scales and on dark energy. ",
        "introduction": "S_d= (V T^d) \\frac {(d-1)!}{2^{d-1}\\pi^{\\frac d2} \\Gamma(d/2)} (d+1) \\zeta(d+1)+ {\\cal O}(V T^d)^0 \\ , \\end{eqnarray} where $T = 1/\\beta$ is the temperature of the chosen Gibbs state. At low temperature/small volume, as for the usual approach after the subtraction of the infinite contribution, thermal entropy goes to zero, although with a different behavior. The only known explicit calculation in standard localization is in 1+1 dimension~\\cite{cala} and in higher dimensions by using AdS/CFT~\\cite{ads}. From those studies it turns out that in standard localization, after the subtraction, thermal entropy is sub-extensive, $S_{\\rm therm}\\simeq (VT^d)^{(d+1)/d}$ at low temperature, whereas our regularized entropy approaches extensivity from above. For small $VT^d$, we obtain \\begin{equation} S\\simeq -VT^d \\ln VT^d + {\\cal O}(V T^d). \\end{equation}       \\subsubsection{NW allows localized states, standard localization does not} \\label{sec:loca}  The usual localization scheme seriously challenges any idea of \\emph{localized state}. In Sec. \\ref{sec:4.1} I recalled that, according to the usual interpretation, a particle instantaneously created at time $t$ at point $\\x$ is described by the state of the theory (eq. \\ref{standardlocalized}) \\begin{equation} \\label{state} |\\x, t \\ {\\rm standard}\\ket  =  \\phi(\\x,t) |0\\ket . \\end{equation} The problem with such a description is that the above state is different from the vacuum at any point $\\x ' \\neq \\x$. In other words, the particle is not localized at all. There is, in fact, a  precise and intuitive sense in which a state is localized. One very natural definition is the following. A state $|\\psi\\ket$ is \\emph{strictly localized} \\cite{strict} outside of $P$ if, for any possible observable $A$ in ${\\cal A}(P)$, $\\bra\\psi|A|\\psi\\ket=\\bra 0|A|0\\ket$. In other words, if we excite some degrees of freedom that are ``strictly localized'' outside $P$, the state of affairs inside $P$ is the same as the one of the vacuum.  It turns out that no state with finite energy has this property in the standard localization scheme. The state \\eqref{state} which is commonly described as ``a particle at position $\\x$'' is in fact different from the vacuum in any region $P$ with $\\x \\notin P$, i.e. $\\rho_P \\equiv {\\rm Tr}_R  |\\x, t \\ {\\rm standard}\\ket \\bra \\x, t \\ {\\rm standard} | \\neq {\\rm Tr}_R  |0\\ket \\bra 0|$. This property, that can be traced back to Reeh-Schlieder theorem \\cite{clift,terno}, is related, once again, with the fact that the vacuum is entangled in the standard scheme. On the other hand, low energy excitations can be ``strictly local'' in the NW scheme because of the factorization \\eqref{identification} that leaves $P$ empty and in its ``local vacuum'' whenever we excite some degrees of freedom somewhere else (i.e. in $R$).  \\subsection{Again on vacuum entanglement vs separability} \\label{sec:separable}  Entanglement is a  well understood -- and measurable -- property of many condensed matter systems in their ground state (see, e.g. \\cite{vidal}). When standard localization is used, the vacuum in quantum quantum field theory shares the same property: different regions of space are entangled, with an entropy proportional to the area of the boundary between them~\\cite{en1,en2,en3,en4,en5}.  One of the most popular interpretations of black hole entropy~\\cite{damour} is in fact to consider it as the entanglement entropy of the quantized fields on the black hole background (e.g. ~\\cite{jac,dvali}). The idea is to trace over the degrees of freedom inside the horizon because they are not accessible from the outside. This gives a UV-divergent answer proportional to the area of the boundary, which is compatible with black hole entropy if the cut-off is suitably chosen. The question that I am going to raise here is whether those correlations that vacuum entropy is accounting for have any operational meaning, i.e., whether they can be used as a source of entanglement in EPR-like experiments. In the black hole case, because of the presence of the event horizon, the correlations between inside and outside are argued to be lost, which should explain the entropy of the black hole itself. But are vacuum correlations ever really accessible anyway, for instance, in flat space? In this section I am going to argue for a negative answer to that question: the ground state correlations are accessible in a condensed matter system, but the vacuum state correlations are not in QFT.    \\subsubsection{Samantha and Vincent}  Consider a quantum spin system with a general nearest-neighbour Hamiltonian. Schematically, \\begin{equation} \\label{ham} H \\, = \\, \\sum_{\\bf j} f(\\sigma_{\\bf j}) + \\sum_{\\bf j} \\sum_{{\\bf k} \\in {\\bf j}} g(\\sigma_{\\bf j} \\otimes \\sigma_{{\\bf k}})\\,  +\\,  \\dots \\ , \\end{equation} where $\\sigma_j$ are the local degrees of freedom, $f$ and $g$ are polynomials, internal indexes and more complicate tensorial structures have been omitted. The symbol $\\in$ here stands for ``is neighbor of\" and the dots possibly stand for next-to-nearest neighbor- or more than two spins- interactions.    The above Hamiltonian aims to resemble  a cut-off field theory, $\\sigma_{\\bf j}$ being the relativistic fields and the nearest neighbors interaction a lattice version of the gradient terms.  Say that there exists a concrete realization of that spin system, and that such an apparatus is physically placed on the table of Samantha, an experimental quantum computer scientist. Samantha is very skilled and has enough instruments in her lab to have full access to the individual spins $\\sigma_j$, she can pick one up and take it out of the system, she can make individual measurements etc\\dots Now let us suppose that the dynamics internal to such a condensed matter system is extremely rich, say that (\\ref{ham}) is a lattice version of the Standard Model or some lower-energy truncation of it. Then, if the system is big enough, there might be something similar to our world within it: some eigenstates of (\\ref{ham}) will correspond to stable elements, more complicate configurations will correspond to scattering experiments. There will be quantum states, evolving according to (\\ref{ham}), that we should be authorized to interpret as ``measurements\" going on within the system. Finally, we will have observers! Call Vincent one of those observers. Vincent is made by the excitations of (\\ref{ham}) as much as we are made of atoms, molecules etc \\dots  It is usually the case that the Hamiltonian (\\ref{ham}) has an entangled ground state. More precisely, the ground state is entangled with respect to the quantum partition defined by the spins $\\sigma_j$, which corresponds to the standard localization. For Samantha such an entanglement has a very clear operational meaning. As we have assumed, she has full access to the spins' degrees of freedom. Therefore, if she can manage to cool the system down and bring it close to its ground state, then she can divide the spins into two groups and measure the Bell correlations between them. Alternatively, she can take few spins out of the chain and see whether the state of the system rearranges afterwards.  For Vincent the situation is different since (\\ref{ham}) is his most fundamental -- field theory -- description. Typical objects in Vincent's world are not localized in Samantha's sense because they are eigenstates or quasi-eigentstates configurations of (\\ref{ham}). This is discussed in Sec. \\ref{sam} below and in~\\cite{fabio2} in more detail. In quantum mechanics we are used to prepare states in arbitrary configurations. However, what seems peculiar of field theory at low density and big volume is that generic states generally evolve by radiating away decay products  until we are left, in a sufficiently large region of space, with field configurations which are approximate eigenstates of the total Hamiltonian: stable particles, molecules, crystals etc\\dots This applies even more dramatically to strongly interacting theories: scattering products are known to hadronize very rapidly if they are not QCD eigenstates.  We are arguing that Vincent does not have access to the spin degrees of freedom and, therefore, to the ground state correlations. The objects in his world (particles, atoms  etc\\dots) are approximate eigenstates-excitations of \\eqref{ham}. Whenever he measures quantum correlations he is doing so among such objects, not among the original spins. For Vincent, the most natural description in terms of subsystems effectively implies a different TPS than the one given by the degrees of freedom $\\sigma$. He will be clever enough to reconstruct the Hamiltonian \\eqref{ham} that successfully reproduces his experiments and note that such an Hamiltonian is \\emph{local} (first neighbor interactions) with respect to some spin basis. But whenever he tries to operationally define detailed spacetime relations between his objects, instruments etc\\dots he is effectively using a different basis. Effectively, Samantha and Vincent have two different fine-grained descriptions of spacetime.  \\subsubsection{Extracting the correlations from the vacuum; particle detectors} \\label{sam}  What I just argued can be proven wrong straight away by an experiment of vacuum entanglement distillation. Such an experiment has in fact been conjectured, featuring two Unruh-DeWitt detectors. Those are objects with local interactions, i.e., Hamiltonians of the type \\begin{equation} \\label{unruh} H_{\\rm Unruh\\ DeWitt} \\ =\\ \\lambda {\\hat o}\\, \\phi(\\x(t),t), \\end{equation} where $\\lambda$ is a dimensionless coupling and $\\hat o$ is some operator internal to the detector. We can take, for simplicity, a two level detector, for which, ${\\hat o} = {\\hat o}_{\\uparrow} + {\\hat o}_{\\downarrow}$, where ${\\hat o}_{\\uparrow} = |E\\ket \\bra0|$ and ${\\hat o}_{\\downarrow} = ({\\hat o}_{\\uparrow})^\\dagger$ and $|0\\ket$ and $|E\\ket$ are the ground and excited states of the detector.   The ``distillation\" of vacuum entanglement proposed in~\\cite{reznik} relies on the \\emph{dark counts} that those detectors are known to undergo in the vacuum~\\cite{rabi}. Two detectors are kept at rest in Minkowski vacuum at a given relative distance and are both switched on and off at space-like separated events and for time intervals $\\Delta t \\ll 1/\\Delta E$, $\\Delta E$ being the energy gap inside the detectors. Through the interaction with the vacuum the detectors get entangled between each other and the above described dark counts should show EPR-like correlations. If ever realized, this experiment would give a precise operational meaning to the entanglement of the vacuum between different regions of space.  The experiment proposed in~\\cite{reznik} highlights a consistency within the standard localization scheme: the Unruh-DeWitt detector is localized in the standard sense and, consistently, ``feels\" the vacuum correlations that link regions of space according to standard localization. In~\\cite{fabio2} we have argued that realistic detectors have a more trivial response on the vacuum and do not perform \\emph{dark counts} of such a fundamental origin. The reason, again, is that a detector in its ground state, and in the absence of quanta to detect, is represented in the full underlying field theory by an eigenstate of the total Hamiltonian. In fact, being an eigenstate of the Hamiltonian seems incompatible~\\cite{fabio2,terno} with having local QFT interactions on the basis of the Reeh-Schlieder theorem~\\cite{reeh}.  In~\\cite{fabio2} we have studied a simple local relativistic field theory with three scalar fields, $\\chi$, $\\eta$ and $\\phi$ with masses $M$, $\\mb$ and $m$ respectively and interaction Lagrangian ${\\cal L}^I(x) \\sim -\\mu\\, \\chi(x)\\eta(x)\\phi(x)$. The coupling $\\mu$ is massive and is assumed to be smaller than the masses of the other fields in order to allow a perturbative treatment. Moreover, $m\\ll M<\\mb$. The theory contains two stable particles (corresponding to the fields $\\chi$ and $\\phi$, of masses $M$ and $m$ respectively), and a meta-stable one ($\\eta$, of mass $\\mb$). We take the particle $\\chi$ to describe our detector in the ground state. Such a particle can capture a light quanta $\\phi$ and form the meta-stable state $\\eta$, corresponding to our excited detector. To the stable particle $\\chi$, as expected, correspond energy eigenstates of the full theory. This is shown explicitly in ~\\cite{fabio2} by using a dressed particle formalism at first order in perturbation theory.  Then, we derived an effective model of detector (at rest) where the stable particle and the quasi-stable configurations correspond to the two internal levels, ``ground\" and ``excited\", of the detector, and the transitions rate of the original field theory are faithfully reproduced. The detector model (``the Frog\",~\\cite{fabio2}) reads: \\begin{equation} \\label{ours} H_{\\rm Frog} = \\frac{\\mu}{M}\\left({\\hat o}_{\\uparrow} \\Phi^+(\\x) + {\\hat o}_{\\downarrow} \\Phi^-(\\x)\\right), \\end{equation} where, again, ${\\hat o}_{\\uparrow}$, ${\\hat o}_{\\downarrow}$ are the raising and lowering operators of the two level detector and the energy gap inside the detector is $\\Delta E = \\Delta M = \\mb - M$. The complex fields $\\Phi^+(\\x) $  and $\\Phi^-(\\x)$ are defined in terms of the some annihilators produced by the ``dressing\" and which are a combination of the degrees of freedom of the three fields we started with.   Although derived from a perfectly local field theory, our effective detector does not couple locally to the field $\\phi$ but only, separately, to the positive and negative frequencies, $\\Phi^-(\\x)$ and $\\Phi^+(\\x)$, of a scalar field $\\Phi(x) \\equiv \\Phi^-(\\x) + \\Phi^+(\\x)$, which is a ``dressed\" mixture of the original field $\\phi$ and the other fields. The effective coupling~\\eqref{ours} is non-local, because reproduces the behavior of a real-dressed, rather than bare, particle in the theory. Our detector \\eqref{ours} does not click in the vacuum, it does click only in the presence of a quantum. As a corollary, our analysis suggests that the experiment envisioned in~\\cite{reznik} is not doable with real particle detectors, and that the entanglement of the vacuum has no operational implications and cannot possibly be used to create testable correlations\\footnote{Correlations seem to be lost also in the presence of a coarse graining~\\cite{massi}.}.    \\subsection{Discussion}  One may argue that standard localization is just intrinsic to the formalism and that going for alternatives would imply to contradict quantum field theory itself. In fact, in some approaches to QFT, the degrees of freedom $\\phi$ and $\\pi$ are defined \\emph{ab initio} as the localized objects of the theory. It is interesting, in this respect, to make a comparison with the more constructive approach e.g. of the beautiful book \\cite{weinbook}. There, the entire QFT formalism is built step by step in order to give account for scattering experiments and decay processes; rather then a fine-grained spacetime description, the aim is to find amplitudes relating the asymptotic states, which are what is observable in scattering experiments. The relativistic fields $\\phi$, as objects of the theory, are introduced at a later stage (Chap. 5) with the purpose of writing -- automatically and easily -- Lorentz invariant interactions. Those are in fact guaranteed if the Hamiltonian is local, in the usual sense, in the fields $\\phi$, $\\pi$.  When I discuss alternative localization schemes I am not calling into question the local form of the Hamiltonian: in the case under consideration considered in this section, for example, the operator (\\ref{ham}) does have the local form \\begin{equation} \\label{local} H =\\frac{1}{2} \\int d^3 x \\left(\\pi^2 + \\nabla \\phi^2 + m^2 \\phi^2\\right) \\end{equation} when written in terms of $\\phi$ and $\\pi$.  The same operator can also be written in terms of the NW-operators by writing the Hamiltonian in terms of creators and annihilators as in \\eqref{sec4:ham} and inverting the two Fourier transforms \\eqref{newton}, % \\begin{equation} \\label{hamnewton} H \\ = \\int  d^3x \\, d^3y \\, K(|{\\vec x} - {\\vec y}|) \\, a^\\dagger({\\vec x}) a({\\vec y}), \\end{equation} % where the function $K$ is a kernel that dies off as $K \\sim e^{- m|{\\vec x} - {\\vec y}|}$ for $|{\\vec x} - {\\vec y}|\\gg m^{-1}$ where $m$ is the mass of the field and $K \\propto |{\\vec x} - {\\vec y}|^{-4}$ in the massless case.      Independently of the localization scheme, \\eqref{local} and \\eqref{hamnewton} is just the same operator that uniquely sets the dynamics among asymptotic states (cross sections, decay rates etc\\dots). What \\emph{is} dependent on the localization scheme is the fine-grained spacetime description. Take, for instance, the question: ``what's the average energy inside this room now?\". According to the usual prescription, that corresponds to a QFT operator of the type \\begin{equation} H_{\\rm room}^{\\rm standard}\\  \\simeq \\ V \\, \\langle T^{00}\\rangle_{\\rm volume\\ average} \\ \\simeq \\ \\frac{1}{2} \\int_{\\x \\in \\rm room} \\!\\!\\!\\!\\!d^3 x \\left(\\pi^2 + \\nabla \\phi^2 + m^2 \\phi^2\\right)\\  , \\end{equation} where the average symbol is a spatial average and $V$ is the volume of the submanifold ``room\". In the NW localization the answer to the same question is given by a \\emph{different} operator. When limited to the subsystem ``room\", (\\ref{hamnewton}) becomes instead \\begin{equation} \\label{nwhamiltonian} H_{\\rm room}^{\\rm NW} \\ \\simeq \\ V \\, \\langle T^{00}\\rangle_{\\rm volume\\ average} \\ \\simeq \\ \\int_{\\x, \\y \\in \\rm room} \\!\\!\\!\\!\\! d^3x \\, d^3y \\, K(|{\\vec x} - {\\vec y}|) \\, a^\\dagger(\\x) a(\\y) . \\end{equation} Alternatively, one can extend one of the two variables $\\x$, $\\y$ to the whole space, thereby considering also the interaction (inside/outside) part of the energy. Note that the operators $H_{\\rm room}^{\\rm standard}$ and $H_{\\rm room}^{\\rm NW}$ coincide in the limit when the size of the room is much bigger that the Compton wavelength $m^{-1}$ of the field. For a  massless field the two localizations converge when the room is much bigger than the inverse temperature $T^{-1}$  of the quantum state considered.  In summary, as long as we are concerned only with the internal dynamics of the fields, all TPSs describe precisely the same state of affairs: we are just considering different -- equally valid -- partitions into subsystems of the field system, not changing its dynamics (cross sections, decay rates etc\\dots). Things may possibly be different when also gravity is taken into account. Eq. \\eqref{nwhamiltonian} suggests a brutally non-local coupling gravity--matter fields and, effectively, a smearing at small scales of the energy content of a region of space. If we stick with the genuine idea that (semiclassical) gravity is really just the geometry of the physical spacetime, then it is  crucial to understand how the matter degrees of freedom feeding into Einstein equations are ``localized'' in the physical spacetime itself. By incorporating alternative localization schemes, semiclassical gravity inherits from the matter fields a certain amount of non-locality (see eq. \\ref{nwhamiltonian}). ",
        "conclusions": "At the beginning of this review I listed few problems affecting the established low-energy theoretical framework for gravity. As opposed to other famous examples in the history of physics, the problems at hand do not seem to be manifest striking inconsistencies and, therefore, they are not guaranteed to be triggering any imminent major revision or breakthrough. However, in parallel with more standard approaches (model building with scalar fields, new -- local or non-local -- terms in the Lagrangian, string-inspired scenarios, extra dimensions etc\\dots) it is also worth examining the very premises of such an established framework, questioning the parts that are believed to be true but still lack of experimental verification, trying substantially new directions.  More radical approaches, when not manifestly inconsistent, are still inevitably tentative, incomplete, hinting to future work.  The directions mentioned in this review have semiclassical gravity as a starting point and are both based on the idea (Sec. \\ref{sec:2}) that a deeper and more general description of spacetime is that of considering it as a quantum system -- containing the matter fields' degrees of freedom. Accordingly, regions of space at a given time are, very generally, quantum subsystems, rather than pieces of a manifold. In section~\\ref{sec3} I explored the possibility that the metric manifold description is in fact just a small distance approximation; while doing this, I attempt to address directly some of the problems of the standard low-energy framework. As a guide for such a modification, I propose to consider an ultra-strong version of the equivalence principle, that I formulate and motivate in Sec. \\ref{sec:3.3}. In section~\\ref{sec4} I reviewed the advantages of considering a different localization scheme, i.e., a different rationale to associate to a region of space its appropriate degrees of freedom. This suggests a fine-grained spacetime description that differs from the standard one at scales comparable with the Compton wavelength of the fields and/or the temperature of the state considered."
    },
    "0910/0910.2704_arXiv.txt": {
        "abstract": "Much of the science that is made possible by multiwavelength redshift surveys requires the use of photometric redshifts.  But as these surveys become more ambitious, and as we seek to perform increasingly accurate measurements, it becomes crucial to take proper account of the photometric redshift uncertainties.  Ideally the uncertainties can be directly measured using a comparison to spectroscopic redshifts, but this may yield misleading results since spectroscopic samples are frequently small and not representative of the parent photometric samples.  We present a simple and powerful empirical method to constrain photometric redshift uncertainties in the absence of spectroscopic redshifts.  Close pairs of galaxies on the sky have a significant probability of being physically associated, and therefore of lying at nearly the same redshift.  The difference in photometric redshifts in close pairs is therefore a measure of the redshift uncertainty.  Some observed close pairs will arise from chance projections along the line of sight, but it is straightforward to perform a statistical correction for this effect. We demonstrate the technique using both simulated data and actual observations, and discuss how its usefulness can be limited by the presence of systematic photometric redshift errors.  Finally, we use this technique to show how photometric redshift accuracy can depend on galaxy type. ",
        "introduction": "\\label{sec:introduction}  Redshift surveys are a major and growing industry in astronomical research.  The use of photometric, as opposed to spectroscopic, redshifts in these surveys makes it possible to study a much larger number of objects for a given amount of telescope time, and to study the faintest sources.  However photometric redshifts are susceptible to larger random and systematic errors, which can propagate into derived quantities; in order to derive meaningful results using photometric redshifts it is crucial to understand their uncertainties. For instance both random and systematic redshift errors can lead to systematic errors in the luminosity and mass functions \\citep{chen03,marchesini07}.  Studies of galaxy clustering can also be strongly affected \\citep{adelberger05,quadri08}.  In both of these cases, it is possible to correct for the systematic errors in derived quantities if the distribution of photometric redshift errors is well-understood, but in practice this is seldom the case.  Surveys that are designed to constrain the cosmological parameters require especially tightly-constrained photometric redshifts, and significant work has gone in to establishing the photometric redshift accuracy and calibration requirements \\citep[e.g.][]{albrecht06,huterer06,mandelbaum08}.  The standard method used to estimate photometric redshift uncertainties is to directly compare the photometric redshifts to the spectroscopic redshifts for some subset of objects.  However spectroscopic samples are frequently not representative of the full photometric samples; at least at $z \\gtrsim 1$, galaxies with high-confidence spectroscopic redshifts are often brighter, bluer, biased toward a specific sub-population (e.g. Lyman Break Galaxies or AGN), or cover a different redshift range than the full photometric sample.  Furthermore, if the parameters used to calculate photometric redshifts are tuned to minimize the differences between the photometric and spectroscopic redshifts, there is little guarantee that these parameters are optimal for the full photometric sample.  The photometric redshift calculation itself also naturally produces an estimate of the photometric redshift uncertainties.  For template-fitting approaches \\citep[e.g.][]{bolzonella00,brammer08}, the uncertainties are derived from the $\\chi^2(z)$ of the template fits.  However, in practice the uncertainties determined in this way (not to mention the photometric redshifts themselves) can depend quite sensitively on the shape and number of templates used.  Similarly, the uncertainties derived when using empirical photometric redshift algrorithms depend on the quality of the training set \\citep[e.g.][]{collister04}.  In this paper we describe a simple empirical method of using close pairs of objects on the sky to estimate the width and shape of the photometric redshift error distribution.  Because galaxies are strongly clustered in real space, there is a high probability that any one galaxy has nearby neighbors.  Therefore, close pairs of objects will have a significant probability of lying at the same redshift, and the differences in photometric redshifts of paired galaxies can then be used to constrain the redshift errors.  We first illustrate the method using a simulated data set, and then show examples using public data.  Below we use the terms \\emph{physical pairs} when referring to objects that are physically-associated with each other (and thus lie at similar redshifts), and \\emph{projected pairs} when referring to objects that lie at different redshifts.  Additionally, to avoid confusion between an object's actual redshift and its photometric redshift, we refer at times to the former quantity as its \\emph{spectroscopic redshift}.  All magnitudes are on the AB system. ",
        "conclusions": "\\label{sec:summary}  The use of photometric, as opposed to spectroscopic, redshifts makes it possible to study a much larger number of objects for a given amount of telescope time.  But photometric redshift errors will propagate through many different types of analyses, and in practice may comprise a significant source of error in derived quantities.  For this reason a realistic estimate of the distribution of photometric redshift errors is necessary.  Obtaining large, representative samples of spectroscopic redshifts with which to directly measure the photometric redshift errors is observationally expensive, and often completely unfeasible.  For this reason it is of great interest to have a method to estimate the size and distribution of such errors that can be applied with limited, or even non-existent, spectroscopic samples.  In this paper we have presented such a method.  It is based on the idea that a close association of two or more galaxies on the sky may represent a true physical association, in which case the objects will lie at nearly the same redshift and the differences between their photometric redshifts constrains the typical errors.  We have described a simple implementation of this idea that makes use of close galaxy pairs, where the best estimate of the true redshift of a pair is taken to be the mean of the photometric redshifts.  We have described how to estimate the photometric redshift error distribution from the difference in photometric redshifts, as well as how to estimate the catastrophic failure rate.  This technique requires applying a statistical correction for pairs that arise from chance projections along the line of sight, and this is easily done by randomizing the galaxy positions and repeating the analysis.  Although in this paper we have focused on the redshift range $0.5 \\lesssim z \\lesssim 2$, the basic technique can be applied at both significantly lower and higher redshifts.  The concept of using angular associations between galaxies to constrain their redshifts is not entirely new.  \\citet{newman08} uses a cross-correlation between a spectroscopic and a photometric sample of galaxies to infer the true redshift distribution of the photometric sample.  \\citet{erben09} uses the cross-correlation between galaxies selected in two disjoint photometric redshift bins to quantify the photometric redshift errors.  \\citet{kovac09} modifies the photometric redshift probability distribution of objects on the basis of the spectroscopic redshifts of nearby objects.  The technique presented in this paper represents a significant step forward because it is simple to implement, the results are easy to interpert, and it can be applied with limited (or even with a complete lack of) spectroscopic information.  As a first application of our method, we have shown that quiescent galaxies will on average have more accurate photometric redshifts than star-forming galaxies in broadband optical/NIR surveys out to at least $z \\sim 2$.  This is because quiescent galaxies have a strong break in their SEDs near 4000\\AA, and if the location of this break can be pinpointed using the observed photometry, the redshift will be tightly constrained.  Star-forming galaxies, on the other hand, have weaker features in their SEDs over the range of observed wavelengths. Differential photometric redshift errors can lead to differential effects in derived quantities, such as luminosity or mass functions, and should be taken into account when comparing such quantities between samples.  A significant limitation of the method presented in this paper arises from systematic errors in the photometric redshifts.  One type of systematic error, in which all photometric redshifts are biased in one direction, will go completely undetected.  However if a particular class of galaxies (e.g. galaxies on the red sequence) is subject to such a bias, whereas another class (in the blue cloud) is not, then this bias will become apparent by looking at cross-pairs (red-blue pairs). Another type of systematic error is when the photometric redshift distribution shows artificial spikes.  This is particularly problematic, as it means that the photometric redshifts of a pair of objects may both be drawn into the spike, leading to a smaller relative redshift difference and an underestimate of the true redshift errors.  In extreme cases, when this type of error is comparable to the random errors, this effect can lead to highly disturbed error distributions (Fig.~\\ref{fig:cosmos_fainter}).  Both of these types of systematic errors can, however, be accounted for by using pairs where one object has a known spectroscopic redshift.  In principle the results from the close-pairs technique may be subject to subtle biases related to the relationship between galaxy properties and local environment.  For instance, if red sequence galaxies preferentially appear in groups, then they will be over-represented in a sample of close pairs.  Similarly, galaxies with boosted star formation due to close interactions may also be over-represented.  On the other hand, such galaxies may be relatively rare, and it is worthwhile to remember that the number of pairs in a sample is a strong function of the sample size itself, growing like $N^2-N$. Another potential problem is that close pairs of objects may have inaccurate photometry due to blending or poor background subtraction, so it is important to apply a sensible lower limit for the pair separations.  Our method of using angular associations of galaxies to constrain both the redshifts and the redshift errors can be applied and extended in various ways.  Particularly intriguing is the possibility of incorporating information from angular associations directly into photometric redshift codes.  A step in this direction has already been taken by \\citet{kovac09}, who modify the photometric redshift probability distributions of objects that have near neighbors with spectroscopic redshifts; here we simply note that this same idea may be extended to neighbors with only photometric information. Regardless of how the ideas discussed in this paper are used in future, the close-pairs technique is straightforward to apply and should prove to be a useful tool in analyzing data from redshift surveys."
    },
    "0910/0910.5312_arXiv.txt": {
        "abstract": "For a robust interpretation of upcoming observations from {\\sc Planck}  and  LHC experiments it is imperative to understand how the inflationary dynamics of a non-minimally coupled Higgs scalar field may affect  the degeneracy of the inflationary observables. \\\\ We constrain the inflationary observables and the Higgs boson mass during observable inflation by fitting the Hubble function and subsequently the Higgs inflationary potential directly to WMAP5+BAO+SNI We obtain for Higgs mass $m_{Higgs}=143.73^{\\,\\,\\,\\,14.97}_{-6.31}$GeV at 95\\% CL for the central value of top quark mass. We show that the inflation driven by a non-minimally coupled scalar field to Einstein gravity leads to significant changes of the inflationary parameters when compared with the similar constraints from the standard inflation. ",
        "introduction": "\\label{S1} The primary goal of particle cosmology is to obtain a concordant description of the early evolution of the Universe consistent with both unified field theory and astrophysical and cosmological measurements. Inflation is the most simple and robust theory able to explain the astrophysical and cosmological observations, providing at the same time self-consistent primordial initial conditions \\cite{Staro80} and mechanisms for the quantum generation of scalar (curvature) and tensor (gravitational waves) perturbations \\cite{Muk81}.\\\\ In the simplest class of inflationary models, inflation is driven by a single scalar field $\\phi$ (inflaton) with some potential $V(\\phi)$ minimally coupled to  Einstein gravity. The perturbations are predicted to be adiabatic, nearly scale-invariant and Gaussian distributed, resulting in an effectively flat Universe.\\\\ The WMAP 5-year CMB measurements alone \\cite{Dunkeley} or complemented with other cosmological datasets \\cite{Komatsu} support  the standard inflationary predictions of a nearly flat Universe with adiabatic initial density perturbations. In particular, the detected anti-correlations between temperature and E-mode polarization anisotropy on degree scales \\cite{Nolta} provide strong evidence for correlation on length scales beyond the Hubble radius.  Alternatively one can look on inflationary dynamics based on the structure of the Standard Model (SM) of particle physics. Recently, Bezrukov and Shaposhnikov \\cite{BeSh1} reported the  possibility that the SM with an additional non-minimally coupled term of the Higgs field and Ricci scalar can give rise to inflation without the need for additional degrees of freedom already present in  SM. The resultant Higgs inflaton  effective  potential in the inflationary domain allows for a range of cosmologically values for the Higgs boson mass even after the inclusion of the radiative corrections \\cite{BaSt1,BaSt2,BeSh,dSW,BGS}. More precisely, the present cosmological constraints on the Higgs mass are based on mapping between the Renormalization Group (RG) flow and the spectral index of the curvature perturbations emerging from the analysis of WMAP 5-year data complemented with astrophysical distance measurements \\cite{Komatsu}.\\\\ However, after the publication of WMAP results, many works devoted to the reconstruction of the inflationary dynamics during observable inflation \\cite{Julien,Popa} show the compatibility between WMAP results and the predictions of the standard inflation driven by a single scalar field with some potential $V(\\phi)$ which is minimally coupled with Einstein gravity.  In order to have a robust interpretation of upcoming observations from {\\sc Planck} \\cite{Planck}  and  LHC \\cite{LHC} experiments it is imperative to understand how the inflationary dynamics of a non-minimally coupled Higgs scalar field may affect  the degeneracy of the inflationary observables. \\\\ In this paper we aim to constrain the inflationary observables and the Higgs boson mass during observable inflation by fitting the Hubble function $H(\\varphi)$, and subsequently the Higgs inflationary potential $V(\\varphi)$, directly to WMAP 5-year data \\cite{Dunkeley,Komatsu} complemented  with geometric probes  from the Type Ia supernovae (SN) distance-redshift relation and the baryon acoustic oscillations (BAO) measurements. ",
        "conclusions": "In order to have a robust interpretation of upcoming observations from {\\sc Planck} \\cite{Planck}  and  LHC \\cite{LHC} experiments it is imperative to understand how the inflationary dynamics of a non-minimally coupled Higgs scalar field may affect  the degeneracy of the inflationary observables. \\\\ We constrain the inflationary observables and the Higgs boson mass during observable inflation by fitting the Hubble function $H(\\varphi)$, and subsequently the Higgs inflationary potential $V(\\varphi)$, directly to WMAP 5-year data complemented  with geometric probes  from the Type Ia supernovae (SN) distance-redshift relation and the baryon acoustic oscillations (BAO) measurements. We obtain a Higgs mass value of $m_{Higgs}=143.73^{\\,\\,\\,\\,14.97}_{-6.31}$GeV at 95\\% CL for the central value of top quark mass $m_{Top}$=171.3 GeV. Our result is compatible with the previous constrains of the Higgs mass from cosmological data \\cite{BaSt1,BaSt2,BeSh,dSW,BGS}, however the upper and lower bounds of the Higgs mass obtained from our analysis are tightly constrained.  The strong correlation between Higgs mass and the Higgs quadratic coupling $\\lambda$ and the larger degeneracy between the Higgs mass and non-minimal coupling constant $\\xi$, indicate the existence of a degeneracy between $\\xi$ and $\\lambda$ that affect the Higgs mass value inferred from cosmological and astrophysical measurements.  We also show that the inflation driven by a non-minimally coupled scalar field to Einstein gravity leads to significant changes of the inflationary parameters when compared with the similar constraints from the standard inflation with minimally coupled scalar field.   \\vspace{1cm} {\\bf Acknowledgments}\\\\ \\\\ The author acknowledge D. Ghilencea for usefull discussions.\\\\ This work  was partially supported by CNCSIS Contract 539/2009.   \\begin{table} \\caption{The mean values and 95\\% CL lower and upper intervals of the derived parameters obtained from the fit of the inflation model with non-minimally coupled Higgs scalar field and the standard inflation model with minimally coupled scalar field to WMAP5+SN+BAO dataset. All parameters are computed at the Hubble radius crossing $k_*$=0.002 Mpc$^{-1}$.} \\begin{center} \\begin{tabular}{lcc} \\hline \\hline \\\\ Model & Standard Inflation   &  Higgs Inflation  \\\\ Parameter &                  &                   \\\\ \\hline \\\\ $\\Omega_bh^2$&$0.022_{0.021}^{0.023}$&$0.023_{0.022}^{0.024}$   \\\\ $\\Omega_ch^2$& $0.111_{0.106}^{0.117}$  &$  0.114_{0.107}^{0.121}  $\\\\ $\\tau$       & $0.082_{0.055}^{0.109}$  &$   0.094_{0.034}^{0.135} $\\\\ $\\theta_s$   & $1.039_{1.034}^{1.045}$  &$ 1.044_{1.037}^{1.048} $   \\\\ ${\\rm ln}[ 10^{10}A^2_S ]$ & $3.143_{3.083}^{3.201}$  & $3.317_{2.999}^{3.252}$   \\\\ $n_S$ &  $0.956_{0.932}^{0.979} $ & $ 0.977_{0.944}^{0.997}         $            \\\\ $\\alpha_S$   &   $-0.012_{-0.038}^{\\,\\,\\,\\,0.021}$ &$0.003_{-0.0021}^{\\,\\,\\,\\,0.0133}$ \\\\ $R$ &     $< 0.556$         & $< 0.234$                 \\\\ \\hline $m_{Higgs}$(GeV)&  &$142.737_{136.431}^{158.707}$\\\\ $\\lambda$   &  &$0.168_{0.153}^{0.185}$\\\\ $\\xi \\times 10^{-5}$& &  $1.871_{0.118}^{0.280}$\\\\ \\hline \\hline \\end{tabular} \\end{center} \\end{table}  \\newpage"
    },
    "0910/0910.5638_arXiv.txt": {
        "abstract": "In this paper we generalize the flux ratio method \\cite{2009A&A...500L..17B} to the case of two luminosity indicators and search the optimal luminosity-flux ratio relations on a set of spectra whose phases are around not only the date of bright light but also other time. With these relations, a new method is proposed to constrain the host galaxy extinction of SN Ia and its distance. It is first applied to the low redshift supernovas and then to the high redshift ones. The results of the low redshift supernovas indicate that the flux ratio method can indeed give well constraint on the host galaxy extinction parameter E(B-V), but weaker constraints on $R_{V}$. The high redshift supernova spectra are processed by the same method as the low redshift ones besides some differences due to their high redshift. Among 16 high redshift supernovas, 15 are fitted very well except 03D1gt. Based on these distances, Hubble diagram is drew and the contents of the Universe are analyzed. It supports an acceleration behavior in the late Universe. Therefore, the flux ratio method can give constraints on the host galaxy extinction and supernova distance independently. We believe, through further studies, it may provide a precise tool to probe the acceleration of the Universe than before. ",
        "introduction": "\\label{sn:introduction} The Type Ia Supernova (SN Ia) as a standard candle plays an important role in the cosmological probes. It is the relations proposed in \\cite{1993ApJ...413L.105P} that make the measurement of decelerator factor of the universe  possible. Based on this relationship, a set of improved techniques are developed, such as the Multicolor Light-Curve Shape method (MLCS) \\cite{1996ApJ...473...88R}, the Stretch method \\cite{1997ApJ...483..565P}, the Spectral Adaptive Light curve Template method (SALT) \\cite{2005A&A...443..781G} and so on. Using these methods, the luminosity distance of a supernova can be measured to a precision $7-10 \\%$. In contrast to these methods which use the light curve shape and color information to standardize the supernova luminosity, some authors use spectral features as the luminosity indicators, for example, \\cite{1995ApJ...455L.147N}, \\cite{2006ApJ...647..513B}, \\cite{2006MNRAS.370..299H}, \\cite{2008A&A...477..717B}. Although these methods are very interesting, till now they are not so much competitive as the above mentioned methods.   The SN Ia luminosity is affected by many factors. One of the most important is the host galaxy extinction. As for this, there exits several methods to estimate it. The first one is the Lira-Phillips relation \\cite{1999AJ....118.1766P}, which show that the B-V color curve of SN Ia with small host galaxy extinction evolves in a similar way within [30-90] days after the maximum light. Therefore, comparing the observed B-V color curve with this color evolution, one can constrain the host galaxy extinction for an individual SN. The second one is the Multicolor Light Curves Shape (MLCS) method \\cite{1996ApJ...473...88R}, which compare the observed light curves with the template light curves. The third one is to measure the equivalent width of the interstellar Na I D doublet which is correlate to the line of sight dust \\cite{1990A&A...237...79B,1997A&A...318..269M}.   Recently, \\cite{2009A&A...500L..17B} searched the optimal correlation between the B maximum M and the flux ratios at any two wavelengths. It is claimed that the scatter can be decreased to 0.12. However, in that paper, the phases of spectra are required within $\\pm 2.5$ days around the day of B maximum $t_0$, and the host galaxy dust extinction are corrected in an indirect method. Here we search the possible relationships between the B maximum M and the flux ratios at different phases, for example, $t = -6 \\sim 14$ days, with all the supernova spectra are deredden to correct for both Galaxy and host galaxy extinction directly. Furthermore, we search the optimal linear relationships by introducing two flux ratios $R_{\\lambda_1 / \\lambda_2},R_{\\lambda_3 / \\lambda_4}$ as the luminosity indicators. It is expected that a part of the intrinsic scatter in the one luminosity indicator case can be compensated by introducing another luminosity indicator, therefore the final scatter can be deceased as much as possible.    After corrected for the host galaxy dust contamination directly, small luminosity scatters of SN Ia can be obtained from the flux ratio method, especially in the case of two flux ratios being used as the luminosity indicators. With this, it is possible to fit the host galaxy extinction and distance modulus at the same time. Such an attempt to constrain the host galaxy extinction and distance modulus is different from the methods employed before. It can provide another independent estimation on host galaxy extinction and distance modulus. Besides this, several fitting results for the same supernova can be compared with each other since the luminosity-flux ratio (L-F) relation has been generalized to other supernova phases besides the maximum day.  The aim of this paper is to develop the flux ratio method and make it to be a reliable tool to estimate the distance of type Ia supernova and its host galaxy extinction. We will go through all the procedures to produce a Hubble diagram from the supernova spectra by using the flux ratio method. Considering that the number of the train spectra is finite and it is the first time to apply the method, we pay more attentions on the method itself than the detailed results.  This paper is organized as follows. Section \\ref{sn:relation} concentrates on the luminosity-flux ratio relation, in which the SN Ia luminosities are fitted with one luminosity indicator or two luminosity indicators respectively and the optimal luminosity scatters are obtained by searching all the wavelength pairs at each day within [-6,14]. In section \\ref{sn:fitting}, a method is proposed to constrain the host galaxy extinction and distance of the supernova. Then it is applied to the low redshift samples. In section \\ref{sn:highz}, the flux ratio method is applied to the high redshift supernova spectra, and the difference between high redshift supernova samples and the low ones are emphasized. At the end of this section, the Hubble diagram and the $(\\Omega_{m}, \\Omega_{\\Lambda})$ contour map are given. Finally we give the discussion and summary in the last section. ",
        "conclusions": "In this method, one of the extinction parameter $E(B-V)$ can be constrained very well, while $R_V$ can be only constrained loosely or even can not be constrained. It is noted that according to the CCM extinction law \\begin{eqnarray} A_{\\lambda_1} - A_{\\lambda_2} & = & A_V[(a_{\\lambda_1}-a_{\\lambda_2}) +(b_{\\lambda_1}-b_{\\lambda_2})/R_V]  \\nonumber\\\\ & = & E(B-V) [R_V (a_{\\lambda_1}-a_{\\lambda_2}) + b_{\\lambda_1}-b_{\\lambda_2}], \\end{eqnarray} In most cases $R_V |(a_{\\lambda_1}-a_{\\lambda_2})|$ is smaller than $|b_{\\lambda_1}-b_{\\lambda_2}|$. Thus, $(A_{\\lambda_1} - A_{\\lambda_2})$ and also the right hand side of Eq. (\\ref{eq2}) or Eq. (\\ref{eq3}) is mainly determined by $E(B-V)(b_{\\lambda_1}-b_{\\lambda_2})$. It also should be noted that $A_B \\approx E(B-V) (R_V +1)$. Thus using a few relations Eq. (\\ref{eq2}) or Eq. (\\ref{eq3}), one can determine $E(B-V)$ very well, but cannot determine $R_V$ or only give a weak constraint on it. Even it can give weak constraint on $R_V$, there exist degeneracy between $R_V$ and $\\mu$, since $(R_V,\\mu)$ appear in the PDF in the form of $E(B-V)R_V + \\mu$ approximately. In such a case, for a given $E(B-V)$, $(R_V,\\mu)$ can only be determined to a region along the line $E(B-V)R_V + \\mu \\approx const$. This can be seen from the contour map of $(R_V,M^*[=M+A_{B,gal}+A_{B,host}=m-k-\\mu])$ for SN 1998dm. However, $R_V$ also affects $(A_{\\lambda_1} - A_{\\lambda_2})$ and also the right hand side of Eq. (\\ref{eq2}) or Eq. (\\ref{eq3}). For different luminosity-flux ratio relations their importance are different. Thus $R_V$ still can be  determined from this method in principle, but for the L-F relations obtained here, the constraint on $R_{V}$ is weak and even not converge (approach to infinity or zero, since in the fitting procedure $R_{V}$ is assumed to be great than zero). In order to break the degeneracy and give more better constraint on host galaxy extinction and distance, one needs more better luminosity-flux ratio relations which are derived from the SN Ia samples that the host galaxy extinction is better corrected. So in the fitting procedure, $R_V$ is always fixed at some constant. This do not affect the constraint on the values of $E(B-V)$ largely.  \\begin{figure}[htbp] \\begin{center} \\includegraphics[width=0.7\\textwidth]{dm.eps} \\end{center} \\caption{Contour map of the effective magnitude $M^*$ of SN 1998dm and the extinction parameter $R_V$ of its host galaxy, the two regions represent $1 \\sigma, 2 \\sigma$ separately.} \\end{figure}   \\subsection{Fitting procedures and results} For a given spectrum with phase $t$, one can fit this spectrum with different L-F relations whose phase is around $t$. Using one L-F realtions one get a result. Among these results, one can choose the result which make the smallest change of the spectrum (due to the host galaxy extinction) as the best one. The above choice means that in the case of large host galaxy extinction the result with smallest $E(B-V)$ is the best one, while in the case of small extinction the result with $E(B-V)$ that approaches zero is the best one. This result is always better than the one obtained by applying the L-F relations at phase $t$. We also attempt to choice the result which has the smallest $\\chi^2 / dof$ as the best one, but it does not work as well as the result which has the smaller change of the spectrum. This point should be checked in a further study.   Since the constraint on $R_V$ is weak, to fit the low redshift supernova spectra listed in Table \\ref{tbl1x}, $R_{V}$ takes values provided there. Thus the fitting parameter is $(E(B-V), \\mu)$. When the apparent magnitude is not known, we choose the fitting parameter as $(E(B-V), M^*=M+A_{B,gal}+A_{B,host}=m-k-\\mu)$ (one can also choose $(E(B-V), M)$). Figure \\ref{dtx} illuminate how $(E(B-V), \\mu)$ converges with the number of L-F relations used in the fitting. Table \\ref{res:lowz} present all the results of low supernova samples. Besides the best ones, the results determined by L-F relations whose phase is given in Table \\ref{tbl1x} are also listed. It is found that for most supernovas, the first set of results is close to the values listed in Table \\ref{tbl1x}, besides SN 1989B whose host galaxy extinction is estimated to be larger than that estimated by \\cite{2007ApJ...659..122J} (correspondingly the distance is smaller than that estimated by \\cite{2007ApJ...659..122J}). Thus the method proposed in this section can indeed give information on the distance and its host galaxy extinction.  It is noted that the distances obtained from the spectra with different phases are generally consistent with each other but not completely. On the one side, this fact reflects that the method is valid, on the other hand it indicates there exist some uncertainty in the relations Eq. (\\ref{eq2}) or Eq. (\\ref{eq3}) which are obtained from a 38 SN Ia samples. This can be regarded as one advantage of the flux ratio method. That is, from the same method, different distances can be obtained to monitor the inconsistency of the method. The smaller is the difference between these distances, the more reliable to apply this method. Furthermore, one can get two different distances by using the same spectrum from luminosity-flux ratio Eq. (\\ref{eq2}) and Eq. (\\ref{eq3}) respectively, it also can reflect how well the F-L relations are. In short, it is possible to see the reliability of the flux ratio method from itself when one have two or more supernova spectrum with different phase.  \\begin{figure}[htbp] \\begin{center} \\includegraphics[width=0.7\\textwidth]{dtx.eps} \\end{center} \\caption{Dependence of supernova distance and its host galaxy extinction of SN 97dt on the number of L-F relations used in the fitting. The upper points represent $(\\mu-30,n)$, while the bottom points denote $(E(B-V),n)$.} \\label{dtx} \\end{figure}   \\begin{deluxetable}{llllllll} \\tabletypesize{\\scriptsize} \\tablecaption{Extinction and distance of low redshift supernovas.} \\tablewidth{0pt} \\tablehead{ \\colhead{SN Ia}   & \\colhead{Date}  & \\colhead{$E(B-V)^1_{gal}$ \\dag}   & \\colhead{$\\mu^1$ or $M^{1*}$ \\ddag}  & \\colhead{phase1 \\$}  & \\colhead{$E(B-V)^0_{gal}$ \\dag}   & \\colhead{$\\mu^2$ or $M^{0*}$ \\ddag}  & \\colhead{phase0 \\$} } \\startdata 1986G  & 1986-05-08 & 0.78(0.00) & 27.42(0.03) & -1& 0.84(0.01) & 27.12(0.04)  & -3\\\\ 1989B  & 1989-02-17 & 0.64(0.01) & 29.12(0.05) & 9 & 0.94(0.01) & 28.19(0.08)  & 10\\\\ 1997dt & 1997-12-04 & 0.42(0.01) & 33.06(0.05) & 1 & 0.42(0.01) & 33.06(0.05)  &  1\\\\ 1998bu & 1998-05-28 & 0.27(0.01) & 30.60(0.03) & 9 & 0.27(0.01) & 30.60(0.03)  &  9 \\\\ 1998bu & 1998-05-29 & 0.23(0.01) & 30.71(0.03) & 9 & 0.31(0.01) & 30.48(0.04)  & 10\\\\ 1998dk & 1998-09-10 & 0.22(0.01) & 18.61(0.02) & 8 & 0.26(0.01) & 18.57(0.04)  & 10\\\\ 1998dk & 1998-09-11 & 0.20(0.01) & 18.67(0.02) & 9 & 0.27(0.01) & 18.48(0.02)  & 11\\\\ 1998dm & 1998-09-10 & 0.27(0.01) & 18.28(0.04) & 4 & 0.31(0.02) & 18.16(0.06)  & 5\\\\ 1998dm & 1998-09-11 & 0.26(0.01) & 18.26(0.04) & 4 & 0.31(0.01) & 18.28(0.04)  & 6\\\\ 1998ec & 1998-09-30 & 0.16(0.01) & 18.51(0.03) & 1 & 0.28(0.00) & 18.21(0.01)  & -1\\\\ 1999cl & 1999-06-14 & 1.11(0.01) & 30.80(0.05) & 1 & 1.11(0.01) & 30.80(0.05)  & 1\\\\ 1999gd & 1999-12-08 & 0.41(0.01) & 34.67(0.04) & 1 & 0.54(0.01) & 34.26(0.06)  &  4\\\\ 2000B  & 2000-01-28 & 0.13(0.01) & 18.78(0.03) & 6 & 0.31(0.01) & 18.28(0.02)  & 9\\\\ 2002bo & 2002-03-23 & 0.37(0.01) & 32.29(0.06) & 1 & 0.54(0.00) & 31.81(0.05)  & -1\\\\ 2002bo & 2002-03-28 & 0.40(0.01) & 32.38(0.06) & 6 & 0.61(0.01) & 31.35(0.06)  & 5\\\\ 2003cg & 2003-03-30 & 1.22(0.01) & 31.83(0.03) & 1 & 1.35(0.00) & 31.61(0.03) & -1 \\enddata  \\tablenotetext{\\dag}{$E(B-V)^{1(0)}_{gal}$ are the host extinction obtained by using the L-F relation at phase1 (phase0).} \\tablenotetext{\\ddag}{$\\mu^{1(0)}$ are distance obtained by using the L-F relation at phase1 (phase0).} \\tablenotetext{\\$}{Phase0 is time of spectrum relative to the B magnitude maximum date listed in Table \\ref{tbl1}, while phase1 is the phase of the L-F relation , at which the smallest change is needed to fit the spectrum.}  \\label{res:lowz} \\end{deluxetable}   \\label{sn:discussion1}   From the above analysis, the distance of supernova and its host galaxy extinction can be probed from the supernova spectrum, which is independent from other method. Actually, most type Ia supernovas has spectroscopic follow-up, so it is possible to estimate their distance and host galaxy extinction from the flux ratio method immediately. Furthermore, when there are two or more spectra of one type Ia supernova, the results can be check with each other. Another advantage of the method is that not one particular L-F relation but many of them play the decisive role in the distance estimation.  Of course, there are a lot of possible factors which affect the flux ratio method. The most important points are: 1) the supernova spectra train data set do not contain some one type of supernova spectrum, then the L-F relations may not fit this type of supernova rightly; 2) the total distortion of the supernova spectrum in the spectrograph, this can be monitored by checking the difference between the observed magnitude and that synthesized from spectra;  3) the host galaxy contamination, lots of supernova spectra have a large fraction of host galaxy component. To exploit these spectra, it should subtract the host galaxy components from them cleanly. Any case mentioned above will leads to wrong or bad results, may be the SNLS supernova 03D1gt is such a case. The other possible factors to affect the method includes: 4) the redshift uncertain , both due to the peculiar velocity and the measurement, since in the procedure of deredshift, the error of redshift will leads some uncertain of the spectra in the rest frame; 5) extinction of Milky Way and host galaxy, the host galaxy extinction for the train spectra will affect the L-F relations and then affect the distance estimation of other supernova, the uncertain in Milky dust extinction also will leads to systematic error in the distance estimation; 6) the noise of the spectrum in the measurement; 7) the uncertain of the phase of the train spectra. All these factors will affect the L-F relation, which make it easy to understand the following things. First, the host galaxy extinction $E(B-V)$ is only an effective parameter which contains not only the effect of host galaxy extinction but also other factors which may change the spectrum. When the host galaxy extinction is large, it will dominate among all the possible factors as found in section \\ref{sn:fitting}. Second, there is no or only weak constraint on $R_V$. It is stated in the previous section that this method depends on $R_V$ loosely even in the ideal case. This dependence will be easily broken by the factors listed above, thus only in a more precise analysis, the flux ratio method can provide an effective constraints on $R_V$.  Finally, let us summary the contents of this paper. First, we generalized the flux ratio method presented in \\cite{2009A&A...500L..17B} to supernova spectrum phase which is not around the date of B magnitude $t_{0}$. This is useful when one wants to use this method to study the universe expansion, since generally speaking there are a lot of supernovas which do not have spectrum around the phase $t_{0}$ but have spectrum at other phase. Second, we search the optimal linear relationships between B maximum magnitude and flux ratios using two flux ratios as the luminosity indicators and obtained a scatter which is small than 0.1. The absolute value of the scatter is already small, but more important is that it will improve the scatter compared with that obtained by using only one flux ratio as the luminosity indicator. This improvement possibly comes from two sides, one is that in some case the intrinsic scatter in one flux ratio case can be offset by introducing another flux ratio, the other is that the scatter due to the spectrum  which is not at the same phase maybe be decreased in this method. In \\cite{2009A&A...500L..17B}, the linear relation using one flux ratio as the luminosity indicator is studied at the phase $t=0$ on a better spectra samples obtained by the Nearby Supernova Factory collaboration \\cite{2002SPIE.4836...61A}, whose scatter can be smaller than $0.12$. The result is already very good, but we believe that after using two flux ratios as luminosity indicators, the scatter can be decreased further to a very low level. Third, we try to correct the host galaxy extinction directly in this paper. \\cite{2009A&A...500L..17B} try to use flux rations which corrected for the host galaxy extinction and intrinsic dispersion without distinguishing the two effects. Here, we have a general host galaxy dust correction, but introduce errors from the uncertainties associated with $A_V$. Since the flux ratio is the ratio of flux at two wavelengths, it is less sensitive to the change of the host galaxy extinction as the flux itself, but the host galaxy dust is still an obstacle to standardize the supernova peak luminosity in the flux ratio method. Without a better host galaxy extinction correction for the supernova samples, one can not get the precise luminosity-flux relation relations. The best method to solve this problem is to redo the analysis in a supernova samples whose dust extinction is small. Fourth, after searching the luminosity-flux ratio relations, we attempt to use it to fit the distances and host galaxy extinctions of low redshift supernovas. The results are consistent with that obtained by other independent methods. However, there exists some degeneracy between distance and the host galaxy extinction, which needs a further study. Last, the same method is applied to the high z supernova samples. Before fitting, we research the L-F relations using a smoother average condition in a small wavelength range. Two kind of distances are estimated, from which we obtained the Hubble diagram and analyzed the contents of the Universe, which support an acceleration Universe. So the flux ratio method can indeed provide another independent way to constrain the distance of SN Ia and its host galaxy extinction."
    },
    "0910/0910.0471_arXiv.txt": {
        "abstract": "We use N-body simulations to study the effects of tides on the kinematical structure of satellite galaxies orbiting a Milky Way-like potential. Our work is motivated by observations of dwarf spheroidal galaxies in the Local Group, for which often a distinction is possible between a cold centrally concentrated metal rich and a hot, extended metal poor population.  Here we focus on the evolution of a spherical satellite with two stellar components set {\\it ab-initio} to be spatially and kinematically segregated, and which are embedded in an extended dark matter halo.  We find that an important attenuation of the initial differences in the distribution of the two stellar components occurs for orbits with small pericentric radii ($r_{per} \\le 20$ kpc).  This is mainly due to: $i)$ the loss of the gravitational support provided by the dark matter component after tidal stripping takes place, which forces a re-configuration of the luminous components, and $ii)$ tides preferentially affect the more extended stellar component, leading to a net decrease in its velocity dispersion as a response for the mass loss, which thus shrinks the kinematical gap.  We apply these ideas to the Sculptor and Carina dwarf spheroidals. Differences in their orbits might help to explain, under the assumption of similar initial configurations, why in the former a clear kinematical separation between metal poor and metal rich stars is apparent, while in Carina this segregation is significantly more subtle. ",
        "introduction": "Dwarf galaxies are by number the most common kind of galaxies in the Local Group. Among all types of dwarf galaxies, such as dwarf irregulars, dwarf ellipticals, transition types and dwarf spheroidal (dSphs), the latter dominate the satellite population of the large spirals.  Dwarf spheroidal galaxies are faint (with luminosities between $10^3-10^7 L_\\odot$), metal poor and tend not to be rotationally supported. The dark matter mass enclosed within the optical extent of these small objects ranges between $10^5-10^8 M_\\odot$ \\citep{walker07}, which extrapolated at large radii would give virial masses of the order of $10^8-10^9 M_\\odot$ and higher \\citep{strigari08}. The star formation histories of these small objects are known to be complex and vary from object to object, but all dSphs contain ancient stellar populations, with ages $>$10 Gyr old.  In most dSphs the ancient stellar component is the dominant one or even the only one present (e.g. Sculptor, Draco, Ursa Minor); there are however a few cases where intermediate age stars (3-6 Gyr old) are also present (such as Carina) or dominate the overall stellar population (e.g. Fornax, Leo~I) \\citep{gallagher94,mateo98,grebel03}.  An intriguing observational feature of dSphs is that, notwithstanding their small sizes, the star formation and chemical enrichment histories are not uniform inside these objects, but varied as exemplified by the presence of spatial gradients in the average metallicity and sometimes age properties of their stellar populations. Harbeck et al. (2001) \\citep{harbeck01} used wide-field imaging to study a set of 9 Local Group dwarfs and examined the spatial distribution of stars in different evolutionary phases as selected from the colour-magnitude diagram (CMD).  In particular, the authors used horizontal branch (HB) stars, known to be ancient ($>10$ Gyr old), and determined the relative spatial distribution of the blue and red (horizontal branch) populations (BHB and RHB, respectively) in order to identify the existence of metallicity/age variations. Harbeck et al. found that segregation of populations may not be uncommon in dSphs although it is not {\\it necessarily} present in all of them.   Further new evidence on the presence of multiple populations came a couple of years later, when Tolstoy et al. (2004) \\citep{tolstoy04} added to wide-field imaging also the information from wide-field spectroscopy, which yielded metallicities and line-of-sight velocities for hundreds individual stars in the Sculptor dSph.  These authors found that metal rich stars in Sculptor are more centrally concentrated and have on average lower velocity dispersions than those of the metal poor component, which on the contrary, show hotter kinematics and are more extended. Clear stellar population gradients have been confirmed for Tucana, Sextans, Sculptor \\& AndVI \\citep{harbeck01}, and more recently in Draco and AndII \\citep{faria07,mc07b}. On the other hand, weak or no gradients have been found in other satellite galaxies of the Local Group as well, such is the case for Carina, Leo I \\citep{koch07}, AndIII \\& AndV$-$VI \\citep{harbeck01}. Carina for example shows a mild segregation between HB and red clump stars, comparable in amplitude to the metallicity gradient found for spectroscopically confirmed Carina red giant stars \\citep{koch06}.  When kinematical information is also available it is found that these stellar populations differences sometimes clearly show up as multiple kinematical components as in the case of Sculptor and Fornax (ref. \\citep{battaglia06}). In other cases such as in Carina the kinematical distinction of stars with different metallicities is more subtle.   The eventual occurrence and evolution of this spatial and kinematical segregation of stars in dwarf galaxies constitutes an interesting problem for models of galaxy formation.  Tolstoy et al (2004) put forward several hypotheses in order to explain their findings in Sculptor: two episodes of star formation with an inactive period between them, the accretion of gas from the dwarf surroundings able to cool and condensate in the core of them or the photo-evaporation of the outer gas layers due to a UV background \\citep{tolstoy04}.  In recent work, Kawata et al. (2006) \\citep{kawata06} followed the formation of a dwarf galaxy in a self$-$consistent cosmological numerical simulation up to redshift $\\sim 6$. These authors found that the complex interplay between star formation and feedback, coupled to chemical evolution, resulted in a system with a marked metallicity gradient. Kawata et al. thus proposed that a sufficiently steep continuous gradient of a {\\it single} stellar population may appear to observations as two kinematically different metal-poor and metal-rich populations.  Even though this constitutes an interesting possibility, the metallicity distribution and velocity dispersion profiles in this model are only marginally consistent with observations, and the system may evolve further (for example in mass) from redshift 6 to the present-day.   Tides might also play an important role on the evolution of multiple stellar components in dSphs. Given the proximity to the primary galaxy, tidal effects are likely to be relevant on satellite galaxies such as the MW dSph. In some models, the strong tidal field induced by the proximity to the primary may lead to fundamental changes in the dwarfs internal configuration. These include the development of bars and bending instabilities that strongly affect the internal kinematics of the satellite, transforming rotationally supported systems into hot spheroidals \\citep{mayer01a,mayer01b}.  Such processes are triggered especially for eccentric orbits, implying that the properties of the dwarfs we observe today around bright galaxies could possibly give us indirect clues about their primordial (original) state, i.e. before interactions with the primary shaped them differently.   In this paper we explore how gravitational effects may induce a mixing of two initially distinct stellar components in dwarf spheroidal galaxies orbiting a MW-like host potential. Our aim is to gauge the evolution of multiple component dwarfs as they move through the host potential and how this depends on their specific orbital path.  In particular, we focus on the differences shown when the same fiducial satellite is placed and evolved on four orbits around the host with different periods and pericenter distances. We therefore do not attempt to explain the origin of such metallicity/kinematical segregation in the dwarfs in the present paper.  We describe the numerical experiments in Section \\S~\\ref{sec:models}, present and discuss our results in Section \\S~\\ref{sec:results} and summarise our findings in Section \\S~\\ref{sec:concl}. ",
        "conclusions": "\\label{sec:concl}  We have studied, by means of N-body numerical simulations, the effects of the tidal forces induced by a host galactic potential on a multi-component satellite galaxy. The model satellite is set up {\\it ad$-$hoc} to initially have two spherical different stellar populations, that are kinematically and spatially segregated.  A more centrally concentrated and with lower velocity dispersion $C_1$ (reminiscent of a ``metal rich'' population) and a more extended, higher velocity (``metal poor'') $C_2$ population. These two components are deeply embedded within a dark matter halo that largely dominates the total mass of the system. This satellite model roughly matches the structural properties of classical Local Group dwarfs. We have followed the evolution of such object in time for $t\\sim 6$ Gyr on four different orbits around a Milky Way-like host.  We find that the ability to distinguish kinematically the two stellar components (set initially in our model to differ by $\\sigma_2-\\sigma_1 \\sim 4$ km/s) is strongly dependent on the amplitude of the tides experienced by the satellite during its orbit.  In cases where a significant amount of mass has been removed, the velocity gap between the more concentrated (colder) stellar population and the more extended (hotter) component can decrease between $\\sim$30-70\\% of its initial value, thus partly erasing the initial kinematical segregation between the stellar populations. The magnitude of this effect is tightly related to the tidal stripping experienced by the satellite, and in particular, the removal of luminous mass is necessary for this effect to be significant.  Such conditions are more easily obtained for orbits restricted to the inner regions of the host potential, whose close pericenter passages and short orbital periods promote the tidal stripping of the satellite's particles.   We apply these ideas to Sculptor and Carina, two of the classical dwarf galaxies of the Milky Way. Sculptor shows a stronger metallicity gradient, together with a lower velocity dispersion for the more centrally concentrated, metal rich stars compared to the metal poor population. On the other hand, in Carina trends are considerably more subtle, with no kinematical distinction between populations detected to date. We argue that differences in the orbital paths of these two dwarfs can be partially responsible for the weakening of the velocity gap in Carina, while leaving Sculptor unaffected due to its more external and longer period orbit.   Further validation of the tidal ``mixing'' scenario that has been proposed in this {\\it Paper} depends upon efforts from theoretical as well as observational studies.  Comprehensive models to understand the formation of the metallicity gradients as well as their likelihood in dwarf galaxies are of fundamental relevance. The structure of the dark matter halos of dwarf spheroidals (core vs cusped) might also play an important role.  Cored profiles could provide less gravitational support to oppose the tidal stripping forces compared to our cusped models, perhaps changing (albeit we expect only quantitatively) the timescales and relevance of the effects studied here.  Larger samples of line of sight velocities for stars associated to dwarfs spheroidals of the Local Group may improve the detectability of double stellar components kinematically segregated in cases where such feature has not yet been identified. Finally, our models predict that some degree of mass depletion on the stellar components must take place in order to considerably affect the kinematical segregation of dwarfs with composite stellar populations. Deep photometric surveys mapping the outskirts of these satellites, specially beyond their tidal radii, should be able to provide definite clues on the existence (or not) of tidally unbound stars associated to each of these objects, a necessary (although not sufficient) condition for our models to apply."
    },
    "0910/0910.3644_arXiv.txt": {
        "abstract": "IceCube is a cubic kilometer neutrino telescope under construction at the South Pole, a successor to the first-generation AMANDA telescope.  IceCube is now three quarters complete, with completion expected in early 2011, and data taken with the partially built detector already provides a sensitivity surpassing the complete AMANDA-II data set.  Results from searches for astrophysical sources of neutrinos and for evidence of dark matter with both AMANDA and IceCube are summarized.  We also discuss plans for Deep Core, an enhancement of IceCube designed to extend its sensitivity to neutrinos below the TeV scale. ",
        "introduction": "Nearly a century after the discovery of the cosmic rays, considerable uncertainty remains regarding the sources accelerating these particles.  It is believed that cosmic rays at the highest energies, above the ``ankle'' in the cosmic ray spectrum around $10^{19}$ eV, are accelerated by extragalactic objects such as gamma ray bursts (GRBs) or active galactic nuclei (AGN). At lower energies, the sources of the cosmic rays are likely galactic objects such as supernova remnants (SNRs), although to date there is no conclusive evidence for cosmic ray acceleration by any of these objects.  The IceCube neutrino telescope, now under construction at the South Pole, attempts to detect very high energy (VHE, roughly TeV to PeV scale) neutrino emission from the accelerators of the cosmic rays. Neutrino emission is expected from such objects, since the highly energetic environments required to accelerate the cosmic rays will likely contain matter or radiation fields with which the accelerated hadrons may interact before escaping the source.  These interactions will produce charged $\\pi$ and $K$ mesons which decay to muon and electron neutrinos (and antineutrinos, which we will not distinguish in these proceedings) \\cite{Learned:2000sw}.  The neutrinos oscillate over distance between their sources and the Earth, leading to a general expectation of flavor equality at detection, although this is not guaranteed: deviations from flavor equality could provide a variety of information regarding the source environment or probe exotic fundamental physics.  At Earth, VHE neutrinos which undergo charged or neutral current interactions with a nucleon can be detected via the Cherenkov radiation emitted by secondary particles produced in the highly inelastic interaction.  Neutrino flavor identification is possible based on the topology of the Cherenkov radiation: $\\nu_\\mu$ charged current interactions produce a highly energetic muon that travels for hundreds of meters or kilometers on a straight trajectory, leading to long straight tracks in the detector.  Electron neutrinos and neutral current interactions of all neutrinos produce localized showers of particles, leading to an approximately spherical emission of Cherenkov light.  Tau neutrinos can produce a variety of signatures, depending on their energy and decay mode \\cite{Learned:1994wg,Beacom:2003nh,DeYoung:2006fg}.  \\begin{figure}[h] \\centering \\includegraphics[width=80mm]{figure1.pdf} \\caption{Schematic of the IceCube detector, including the initial AMANDA array, the new Deep Core low energy extension, and the IceTop air shower array.  The portion of the detector installed as of 2009 is shown in blue, and the Eiffel Tower is shown for scale.} \\label{fig:layout} \\end{figure}  The IceCube telescope uses a cubic kilometer of the Antarctic ice cap as the Cherenkov medium to detect neutrinos, instrumenting the ice with a three dimensional array of photomultiplier tubes (PMTs) housed in digital optical modules (DOMs).  The completed detector will consist of 5160 DOMs attached to 86 vertical strings at depths between 1450 m and 2450 m below the surface of the ice cap.  In addition, an extensive air shower array on the surface, known as IceTop, will detect cosmic rays.  The layout of the array is shown in Fig.~\\ref{fig:layout}.  At present, 59 strings are installed and operational, with the remaining 27 strings to be deployed in the coming two austral summers.  Also shown in Fig.~\\ref{fig:layout} is the AMANDA telescope, which operated as an independent instrument from 2000-2006 and was integrated into IceCube from 2007-2008.  AMANDA consisted of 687 optical modules on 19 strings, with the bulk of the detector at depths of 1500 to 2000 m. The AMANDA optical modules were simpler than their IceCube counterparts, with considerably smaller dynamic range and less ability to resolve individual photoelectrons.  In 2009, AMANDA was decommissioned and the first string of its replacement, Deep Core, was deployed. ",
        "conclusions": ""
    },
    "0910/0910.1707_arXiv.txt": {
        "abstract": "I report on the development of new code to support the Nasmyth and E-W antenna mount types in AIPS which will allow polarisation analysis of observations made using these uncommon antenna configurations. These mount types will probably become more widely spread as they have several advantages, particularly for geodetic observatories. Multi-band observations, with multiple receivers, can only be fitted into telescopes with Nasmyth feeds. These are the requirements for the new generation of geodetic arrays as discussed in IVS2010. Further more the next generation of antennae will also be required to have very high slew rates, and these can be achieved with the E-W mount.  The mount type affects the differential phase between the left and the right hand circular polarisations (LHC and RHC) for different points on the sky. The target antennae for the project is the Yebes 40m telescope, but as that is still under construction the data used as a demonstration was from the Pico Veleta antenna as part of the Global Millimeter VLBI Array (GMVA). For the E-W mount type there are suitable data from the Australian LBA array. I have demonstrated the effectiveness of the changes made and that the Nasmyth and E-W corrections can be applied.  This report does not cover the basics of interferometry, nor regular polarisation calibration. For references on these topics see, for example, Thompson, Moran, and Swenson (2001) and Aaron (1997). ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.3702_arXiv.txt": {
        "abstract": "Modern supernova (SN) surveys are now uncovering stellar explosions at rates that far surpass what the world's spectroscopic resources can handle.  In order to make full use of these SN datasets, it is necessary to use analysis methods that depend only on the survey photometry.  This paper presents two methods for utilizing a set of SN light curve templates to classify SN objects. In the first case we present an updated version of the Bayesian Adaptive Template Matching program (BATM).  To address some shortcomings of that strictly Bayesian approach, we introduce a method for Supernova Ontology with Fuzzy Templates (SOFT), which utilizes Fuzzy Set Theory for the definition and combination of SN light curve models. For well-sampled light curves with a modest signal to noise ratio (S/N$> 10$), the SOFT method can correctly separate thermonuclear (Type Ia) SNe from core collapse SNe with $\\geq$98\\% accuracy. In addition, the SOFT method has the potential to classify supernovae into sub-types, providing photometric identification of very rare or peculiar explosions. The accuracy and precision of the SOFT method is verified using Monte Carlo simulations as well as real SN light curves from the Sloan Digital Sky Survey and the SuperNova Legacy Survey.  In a subsequent paper the SOFT method is extended to address the problem of parameter estimation, providing estimates of redshift, distance, and host galaxy extinction without any spectroscopy. ",
        "introduction": "\\label{sec:ModelParameters}  When we observe a SN light curve, we are collecting information about the flux as a function of wavelength and time, $f_{obs}(\\lambda,t)$.  If we were to make a list of all the relevant parameters that influence the observed flux, we could divide them up into a set of ``physical parameters'' \\bfPhi\\ and a set of ``location parameters'' \\bftheta.  For \\TNSNe the physical parameter set would include things such as the total mass of ejecta , the mass of $^{56}$Ni, the metallicity of the progenitor, the type of companion star, the asphericity of the explosion, etc.:  $$ {\\bf \\Phi} = (M_{ej}, M_{Ni}, Z, ...) $$  \\noindent The \\bfPhi\\ vector could be broadened to include \\CCSNe as well by including characteristics such as the type of stellar explosion, the presence of Hydrogen or Helium envelopes, etc. If we had a complete model for all SN progenitor systems and a perfect SN explosion simulator, then we could define a translation function $T$ that takes as input a parameter vector \\bfPhi\\ and produces a function describing the intrinsic SN luminosity for any wavelength and time:  \\begin{equation} L(\\lambda,t) = T({\\bf \\Phi},\\lambda,t) \\label{eqn:L(lambda,t)} \\end{equation}   For both TN and CC SNe, the object's location in time and space distorts this intrinsic luminosity in a predictable way, and we observe the flux.  This distortion is predictable because we have reasonably good models describing the way SN light is affected by cosmological distances, dust, and time, so we can write the observable flux as:  \\begin{equation} f(\\lambda,t) = \\frac{1}{4\\pi d_L^2} ~ L\\left(\\frac{\\lambda}{1+z},t-t_{pk}\\right) ~ 10^{-0.4 A_{\\lambda/(1+z)}} \\label{eqn:f(lambda,t)} \\end{equation}  \\noindent where z is the redshift, $d_L$ is the luminosity distance, $A_{\\lambda/(1+z)}$ describes the host galaxy extinction, and \\tpk\\ is the time of the light curve peak. Hereafter, we will drop the $\\lambda$ dependency when discussing fluxes, under the requirement that flux comparisons are always to be done with fluxes from the same broad-band filter. Throughout this paper, we use a simplified formulation of the four ``location parameters:''  $$ \\bftheta = (z, \\mue, \\Av, \\tpk) $$  \\noindent The first parameter is the redshift of the object (z), which determines what portion of the spectral energy distribution (SED) is observed in each bandpass, as well as setting the time-scale of the light curve due to cosmological time dilation. For the second parameter we have recast the luminosity distance using the parameter \\mue.  We define \\mue\\ as the difference between the true distance modulus at a given redshift and the value expected for an empty universe at that z:\\footnote{ Note that this formulation for the luminosity distance does not mean that BATM and SOFT are dependent on an assumed cosmology. In Equation~\\ref{eqn:mue} the distance parameter, \\mue, quantifies the dimming (or brightening) of a SN relative to how it would appear in an empty universe.  The only cosmological parameter appearing here is $H_0$, which factors out as a constant term, independent of redshift z. Therefore, when comparing the BATM/SOFT distance estimates of two SNe at different redshifts, there is no bias introduced by assuming a cosmology, so long as we restrict ourselves to relative distance estimates. }  \\begin{equation} \\mu_e = (m-M) - 5 \\log_{10} \\left(\\frac{cz}{H_0}\\left(1+\\frac{z}{2}\\right)\\right) - 25 \\label{eqn:mue} \\end{equation}  \\noindent The third parameter is the time of the light curve peak (\\tpk), which simply shifts the light curve along the time axis. Our final parameter represents the host galaxy extinction (\\Av), and is discussed in detail in \\S\\ref{sec:HostGalaxyExtinction}.  The parameter ``\\Av'' here is a proxy for a potentially complex extinction function that may vary with time and wavelength.  Throughout this paper, we treat \\bftheta\\ as a set of nuisance parameters, to be marginalized over for the purpose of getting a broad classification into one of the established SN categories.   The task of extracting estimates of some parameters of interest (i.e. measuring z and \\mue\\ for cosmology) is addressed in the companion paper.  \\subsection{Light Curve Shape}  The core function of any SN light curve evaluation program is to predict the observable flux from various SN types, over an (infinite) range of possible location parameters \\bftheta.  By comparing the model to observations we can make statements of probability regarding a candidate SN's class and location.  So how does one define the model for predicting light curve shape?  The fundamental problem is that there is a hidden space of physical parameters \\bfPhi\\ that determine SN light curve shape, and we do not have a theoretical model that can provide the translation of \\bfPhi\\ into luminosity as in Eq.~\\ref{eqn:L(lambda,t)}. In parameterized SN light curve fitters such as MLCS, SALT, and SiFTO, this problem is addressed by making the assumption that a set of shape parameters $\\Delta$ can be used as a proxy for the unobservable \\bfPhi\\ vector.\\footnote{ Here $\\Delta$ represents a vector including a width parameter such as $\\Delta m_{15}$ or ``stretch,''  and possibly also a color parameter and other higher order terms.} Parametric fitters thus produce a single model $M(\\Delta)$, which they assert is capable of representing any possible physical parameter set \\bfPhi\\ by varying the parameter(s) $\\Delta$ to modify the shape of the light curve. In general, the $M(\\Delta)$ model is restricted to TN\\,SNe, although in principle a model for CC\\,SNe could be applied, although it would be more complex and less informative.  For BATM and SOFT, we take an alternative approach, in which we do not parameterize the light curve shape in order to represent all (TN)SNe with a single model.  Instead, we use a finite set of discrete models, each one representing the shape of a single well-observed SN from the local universe.   Whereas the parametric approach essentially tries to fit an empirical function across the hidden \\bfPhi\\ space, with BATM we are sampling the \\bfPhi\\ space at discrete locations. ",
        "conclusions": "\\label{sec:Conclusions}  We have presented in detail the updated BATM (Bayesian Adaptive Template Matching) program, and showed that a strictly probabilistic framework is insufficient for SN classification using fixed-shape templates.  To overcome the practical and theoretical shortcomings of the BATM approach we have introduced SOFT, a method based on Fuzzy Set Theory.  This technique allows us to classify SNe into the broad categories of thermonuclear (TN) and core-collapse (CC) SNe using photometry alone.  The SOFT method is demonstrated on simulated observations of synthetic SNe modeled after the Pan-STARRS 1 (PS1) survey.  Sampling a wide range of redshift, distance and extinction we find that SOFT can produce average classification accuracies above 90\\% for \\TNSNe\\ and above 80\\% for \\CCSNe.  Additional verification of the SOFT classification accuracy is presented with 146 \\TNSN\\ light curves from the SDSS-II survey, which again yield better than 90\\% classification accuracy.  Almost all of the SDSS SN misclassifications could be influenced by peculiarities and ambiguities in the light curves themselves.  We suggest two additional non-SN models that can be helpful in identifying anomalous objects within a SN candidate set.  The Zero Flux Model catches SN impostors that have a very low signal to noise ratio, and the Something Else Model provides a flexible, non-physical model to assist in the identification of unusual objects like the peculiar transient SCP 06F6.  A number of further refinements to the SOFT program could improve its functionality and accuracy for automated SN classification.  The strength of any template matching software is limited by the size and quality of the template library it draws from. The set of templates used for this work has been much enhanced in this regard since the earliest incarnation of BATM, but there is much room for further development.  In particular, collecting template light curves with more complete UBVRI data -- especially in the rise time of the light curve --  would allow for more accurate spline fits to define the primary templates. This is particularly important for the CC SNe, for which the available collection of well-sampled light curves and early-time spectra is notably thin.  Adding to the template library some light curves observed at high redshift could provide a useful supplement to the scarce UV data from the low-z sample.  The SOFT method introduces two new parameters that characterize the template library itself: $\\sigma'_j$ quantifies the amount of ``fuzziness'' and the $q$ parameter controls the strength of the Fuzzy combination operator.  These parameters reflect the number and diversity of templates in the template library, and appropriate values can be estimated from a comparison of templates against each other.  We performed a preliminary calibration of these parameters and settled on $\\sigma_{TN}=0.1$ for \\TNSN\\ templates, $\\sigma_{CC}=0.15$ for \\CCSN\\ templates, and q=0.2 for all combinations.   A more precise determination of $\\sigma'_j$ and q using a large training set of well-observed SNe would be a useful next step.  In addition to classifying SNe, the SOFT method can be used to determine parameter estimates for candidate SNe, based on the comparison of light curves. This extension of the SOFT technique is presented in a companion paper (Rodney and Tonry, 2009), where we focus on the use of SOFT to derive redshift and distance estimates from \\TNSNe\\ for constraining cosmological models.   {\\bf Acknowledgments:} We thank the anonymous referee for constructive comments that led to significant improvements in this work. We thank Stephane Blondin for providing access and assistance with the collection of SNID spectra and Saurabh Jha for sharing low-z light curves.  We thank Brian Connolly for helpful discussion and review of an early draft, and Kathleen Robertson for assistance with reference materials.   This work made use of the SUSPECT database of SN light curves and spectra, maintained by Jerod Parrent (\\url{http://bruford.nhn.ou.edu/~suspect})."
    },
    "0910/0910.2619_arXiv.txt": {
        "abstract": "We present a spectroscopic library of late spectral type  stellar templates in the near-IR range 2.15--2.42\\mic, at R=5300--5900 resolution, oriented to support stellar kinematics studies in external galaxies, such as the direct determination of the masses of supermassive black-holes in nearby active (or non-active) galaxies. The combination of high spectral resolution and state-of-the-art instrumentation available in 8-m class telescopes has made the analysis of circumnuclear stellar kinematics using the near-IR CO band heads one of the most used techniques for such studies, and this library aims to provide the supporting datasets required by the higher spectral resolution and larger spectral coverage currently achieved with modern near-IR spectrographs. Examples of the application for kinematical analysis are given for data obtained with two Gemini instruments, but the templates can be easily adjusted for use with other near-IR spectrographs at similar or lower resolution. The example datasets are also used to revisit the ``template mismatch'' effect and the dependence of the velocity dispersion values obtained from the fitting process with the characteristics of the stellar templates. The library is available in electronic form from the Gemini web pages at \\url{http://www.gemini.edu/sciops/instruments/nearir-resources/?q=node/10167} ",
        "introduction": "Spectral templates, usually late type stars,  are required for the analysis of kinematical data on external galaxies or other stellar ensembles \\citep[e.g.][]{emsellem01,marquez03,barbosa06,ganda06,cappellari07,dumas07,riffel08,cappellari09,riffel09}. In the near-infrared (hereafter near-IR), the most commonly used features are the CO overtone bands at $\\lambda > 2.29$\\mic. These band heads originate from evolved stars and are the strongest features in the 1--3\\mic\\ spectral range of stellar systems. The features are deep and sharp, and at least the first two overtones are located in regions of the IR spectrum relatively clean from telluric lines.  Although observational and theoretical libraries exist at lower spectral resolutions \\citep[$R\\leq 3000$, e.g.][]{kleinmann86,wallace97,ramirez97,forster00,ivanov04}, no comprehensive set of stellar kinematic templates was available to be used with the $R\\sim 6000$ configuration of the two Gemini NIR instruments used for stellar population kinematic studies in external galaxies - the Near-infrared Integral Field Spectrograph (NIFS) and the Gemini Near-Infrared Spectrograph (GNIRS) with the 111 l/mm grating  (both long-slit and Integral Field Unit - IFU), and, therefore, all observing programs using those configurations would invariably spend some science time taking a small set of stellar spectra to use as templates. This led to a constant multiplication of data taking, since those targets were defined as program calibrations and were not made available to other users until the end of the default 18-months proprietary period.  During semester 2006B at Gemini South, given the unusually poor conditions over the whole semester, a Director's Discretionary ``poor weather''  program was specifically carried out to provide the NIR community with a larger set of late (F7 to M3 types) stellar spectra in the range 2.24--2.43\\mic, including the four CO overtone bands, at $R \\sim 5900$ resolution.  Most of the targets were also observed at a slightly bluer spectral range (2.15--2.33\\mic) to improve usefulness for NIFS users, overlapping with the red setting on the first two CO bands.   To the original sample of 29 stars observed with GNIRS, another 11 were added from NIFS observations obtained as part of programs GN-2006A-SV-123, GN-2006B-Q-107, and GN-2007A-Q-25, covering the full range 2.1--2.5\\mic\\ at a similar resolution to that of the GNIRS data. ",
        "conclusions": ""
    },
    "0910/0910.1089_arXiv.txt": {
        "abstract": "Stars do not form continuously distributed over star forming galaxies. They form in star clusters of different masses. This nature of clustered star formation is taken into account in the theory of the integrated galactic stellar initial mass function (IGIMF) in which the galaxy-wide IMF (the IGIMF) is calculated by adding all IMFs of young star clusters. For massive stars the IGIMF is steeper than the universal IMF in star clusters and steepens with decreasing SFR which is called the IGIMF-effect. The current SFR and the total H$\\alpha$ luminosity of galaxies therefore scale non-linearly in the IGIMF theory compared to the classical case in which the galaxy-wide IMF is assumed to be constant and identical to the IMF in star clusters. We here apply for the first time the revised SFR-$L_\\mathrm{H\\alpha}$ relation on a sample of local volume star forming galaxies with measured H$\\alpha$ luminosities. The fundamental results are: i) the SFRs of galaxies scale linearly with the total galaxy neutral gas mass, ii) the gas depletion time scales of dwarf irregular and large disk galaxies are about 3~Gyr implying that dwarf galaxies do not have lower star formation efficiencies than large disk galaxies, and iii) the stellar mass buildup times of dwarf and large galaxies are only in agreement with downsizing in the IGIMF context, but contradict downsizing within the  traditional framework that assumes a constant galaxy-wide IMF. ",
        "introduction": "The determination of accurate current star formation rates (SFRs) is  of fundamental importance for the understanding and investigation of the processes relevant for star formation  and galaxy evolution. A commonly used tracer to calculate the current SFR of galaxies is the measurement of the H$\\alpha$ emission line which is linked to the presence of short-lived massive stars.  The standard way to construct relations between the total H$\\alpha$ luminosity and the total SFR of a galaxy is to apply on galaxy-wide scales an invariant stellar initial mass function (IMF) which is determined in star clusters \\citep[eg.][]{kennicutt1983a,gallagher1984a,kennicutt1994a,kennicutt1998b,kennicutt1998a}. This method provides a linear SFR-$L_\\mathrm{H\\alpha}$ relation as long as the assumed galaxy-wide IMF is treated to be constant and independent of the SFR.  A basic result of SFR studies is that $SFR/M_\\mathrm{gas}$ decreases with decreasing total neutral gas mass, $M_\\mathrm{gas}$. Low-gas-mass (dwarf irregular) galaxies turn their gas into stars over much longer time scales  than  high-gas-mass (large disk) galaxies \\citep[eg.][]{skillman2003a,karachentsev2007a,kaisin2008a}. That dwarf galaxies have lower star-formation efficiencies, defined as $\\tau_\\mathrm{gas}^{-1}=SFR/M_\\mathrm{gas}$, than massive galaxies has been a generally accepted fact at the base of most theoretical work on galaxy evolution.  But in the last years it has been shown that the application of the IMF in star clusters to galaxy-wide scales is doubtful \\citep{weidner2003a,weidner2005a,weidner2006a,hoversten2008a,meurer2009a,lee2009a}. Instead, the galaxy-wide IMF has to be calculated by adding all IMFs of all young star clusters leading to an integrated galactic stellar initial mass function (IGIMF). The fundamental property of the IGIMF-theory is that for massive stars it steepens with decreasing total SFR although the shape of the underlying IMF in each star cluster is universal and constant. This is a direct consequence of (i) the star formation process taking place in individual embedded star clusters on the scale of a pc, many of which dissolve when emerging from their molecular cloud cores,  rather than being uniformly distributed over the galaxy and of (ii) low-star-formation regions being  dominated by low-mass star clusters which are void of massive stars \\citep*{weidner2003a,weidner2005a,weidner2006a,weidner2009a}.  The main aim of this paper is not to present a high precission data analysis of galaxies but to demonstrate the change of fundamental galaxy properties when switching from the constant IMF assumption to the IGIMF theory.  We recap the basics of the IGIMF theory and update its observational support in Section~\\ref{sec_igimf}. The required data of the sample of star forming galaxies are tabulated in Section~\\ref{sec_data}.  We then calculate the SFRs of galaxies (Section~\\ref{sec_sfr}), gas depletion time scales (Section~\\ref{sec_gas_depl}), and stellar mass build-up times (Section~\\ref{sec_buildup_times}) using an IGIMF based SFR-$L_\\mathrm{H\\alpha}$ conversion and show the dramatic differences that arise form using a constant-IMF based SFR-$L_\\mathrm{H\\alpha}$ relation. ",
        "conclusions": "We have applied the revised SFR-$L_\\mathrm{H\\alpha}$ relation, which takes the nature of clustered star formation into account, to a significantly larger sample of star forming galaxies of the local volume than was available in \\citet{pflamm-altenburg2007d}.  The SFRs of galaxies and the corresponding gas depletion and stellar-mass buildup time scales are calculated. Comparing these new values with the classical ones which are derived from SFRs based on an assumed constant galaxy-wide IMF, the changes are dramatic:  The SFRs of galaxies scale linearly with the total galaxy neutral gas mass and the corresponding  gas depletion time scales are independent of the galaxy neutral gas mass. This implies that dwarf galaxies have the same star formation efficiencies as large disk galaxies. Furthermore, the stellar-mass buildup times are only compatible with downsizing in the IGIMF context. They are inconsistent with downsizing in the classical theory which assumes an invariant galaxy-wide IMF.  These results follow from the SFR$_\\mathrm{H\\alpha}$/SFR$_\\mathrm{FUV}$ data compiled by \\citet{lee2009a}, independently of the well developed IGIMF theory. The IGIMF theory merely explains the \\citet{lee2009a} results within an empirically well established star-formation framework. The success of the IGIMF theory, developed by \\citet{weidner2003a} and \\citet{weidner2005a,weidner2006a}, is that the \\citet{lee2009a} results have been predicted \\citep{pflamm-altenburg2007d} rather than being adjusted afterwards. The power of the IGIMF theory lies in that it is readily computable and deterministic.  We emphasise that the galaxy properties derived in the IGIMF theory are qualitatively independent of the main parameter of the IGIMF, the slope of the ECMF. For the full range of possible slopes of the ECMF in the IGIMF theory SFRs of dwarf galaxies are significantly higher than derived in the classical paradigm assuming a constant galaxy-wide IMF independent of the global SFR.  The revision of the SFRs of dwarf galaxies suggested here now poses a major challenge for our theoretical understanding of star formation in galaxies which had been developed with the aim of explaining the low star formation efficiencies of dwarf galaxies. It follows that galaxy and cosmic evolution models need a substantial revision requiring further studies."
    },
    "0910/0910.4386_arXiv.txt": {
        "abstract": "To explain important properties of extrasolar planetary systems (eg. close-in hot Jupiters, resonant planets) an evolutionary scenario which allows for radial migration of planets in disks is required. During their formation protoplanets undergo a phase in which they are embedded in the disk and interact gravitationally with it. This planet-disk interaction results in torques (through gravitational forces) acting on the planet that will change its angular momentum and result in a radial migration of the planet through the disk. To determine the outcome of this very important process for planet formation, dedicated high resolution numerical modeling is required. This contribution focusses on some important aspects of the numerical approach that we found essential for obtaining  successful results. We specifically mention the treatment of Coriolis forces, Cartesian grids, and the {\\tt FARGO} method. ",
        "introduction": "After the discovery of now over 300 extrasolar planets (for an always up to date list, see: {\\tt http://exoplanet.eu/} by Jean Schneider) the process of planet formation has now become again a major area of modern astrophysical research. The dynamical properties of the newly discovered planetary system show distinct differences with our own Solar System, see \\citet{2007prpl.conf..685U} for a review. Most noticeable are very close in massive planets (hot Jupiters and Neptunes), high orbital eccentricities, and the occurrence of low order (2:1) mean-motion resonances. Some of these properties require an orbital evolution of the embedded planets through the disk. This migration process is very generic for young planets that are still embedded in the disk \\citep{1980ApJ...241..425G,1986ApJ...309..846L,1997Icar..126..261W}.  The formation of massive planets can take place primarily along two routes: ({\\it i)} through {\\it gravitational instability} of the protoplanetary disk \\citep{1998ApJ...503..923B}. If the disk is sufficiently massive (about $0.1 M_*$) spiral density arms will form which may turn gravitationally unstable resulting in local fragmentation and the formation of high density 'blobs', i.e. the protoplanets. The main advantage of this scenario is its speed, as it is possible to form Jupiter sized planets within 100 orbital periods. A possible drawback is the requirement of relatively rapid cooling mechanisms in the disk, since the fragments can only collapse if they can cool sufficiently rapidly. Also, the existence of massive solid cores in the centers of the planets cannot easily be explained. ({\\it ii)} through the {\\it core accretion} process \\citep{1980PThPh..64..544M}. In this second scenario the early growth of planets is accomplished by coagulation of small dust particles which are embedded in the gaseous protoplanetary disk. They collide, stick together and grow in mass until eventually (after several stages) a solid core of a few earth masses has be assembled. At this stage, gas accretion sets in and the planet can grow in mass up to several Jupiter masses.  Both models of planet formation must take place within a gaseous environment, i.e. the planet is still embedded within the surrounding protoplanetary accretion disk. In this case, the presence of the growing protoplanet in the disk will generate non-axisymmetric disturbances, the spiral arms (see Fig.~\\ref{kley:fig-crida}). These pull gravitationally on the planet or, phrased differently, exert a torque $\\vc{T}$ on the planet. As a consequence the angular momentum $\\vc{L}$ of the planet will be changed $\\dot{\\vc{L}} = \\vc{T}$, and since for circular orbits $\\vc{L}$ is only a function of the semi-major axis $a$ of the planet, the planet has to migrate. Hence, in both planet formation scenarios, the planet will move radially through the disk. Two observational facts are typically taken as indirect evidence that such a planetary migration process has indeed occurred. First, the existence of hot planets close to the star, as in-situ formation is difficult, and second the occurrence of mean motion resonances, as they require special locations of the planets which are very unlikely to occur just by chance. \\begin{figure}[ht] \\centerline{ \\resizebox{0.65\\linewidth}{!}{% \\includegraphics{crida-sat.jpg}} } \\caption{ Density structure of an embedded Saturn mass planet in a protoplanetary disk. Clearly visible are the two spiral arms (red) and the lower density (gap, dark blue) at the orbit of the planet. (Courtesy: Aur{\\'e}lien Crida) } \\label{kley:fig-crida} \\end{figure} To calculate the evolution, i.e. the mass growth and migration of young planets in the disk, numerical simulations are typically employed. In this contribution we focus on the numerical aspect of planet-disk modeling, and will outline some of the necessary requirements to perform those simulations successfully. ",
        "conclusions": ""
    },
    "0910/0910.1330_arXiv.txt": {
        "abstract": "Automated search for star clusters in $J,H,K_s$ data from 2MASS catalog has been performed using the method developed by Koposov et. al (2008). We have found and verified 153 new clusters in the interval of the galactic latitude $-24^{\\circ}<b<24^{\\circ}$. Color excesses $E(B-V)$, distance moduli and ages were determined for 130 new and 14 yet-unstudied known clusters. In this paper, we publish a catalog of coordinates, diameters, and main parameters of all the clusters under study. A special web-site available at \\url{http://ocl.sai.msu.ru} has been developed to facilitate dissemination and scientific usage of the results. ",
        "introduction": "The search for star clusters in the Galaxy is of a great interest for investigators in the recent years. On the one hand, the usage of new methods and instruments in observations allows scientists to involve clusters in solving a larger number of astrophysical problems and raises many new ones, such as: the processes of formation and evolution of young massive clusters, the nature of nuclear star clusters, the presence of multiple populations in globular clusters. On the other hand, the availability of new multiband all-sky surveys stimulates us to search them for new clusters and build a homogeneous catalog of parameters of both newly discovered and already known clusters.  In Paper I (Koposov et al. 2008), we described the new method of an automated search for star clusters as a density peaks in huge stellar catalogs. It is based on the convolution of density maps with a special 2-D filter, which is the difference between two 2-D Gaussian profiles and has zero integral. If convolved with this special filter, the areas of flat or slowly changing background would produce zero signal, whereas the areas of star concentrations would exhibit a high signal. Using this method, we analyzed the distribution of stars in Two Micron All Sky Survey (2MASS) in the field of $16\\times16$ degrees towards the Galactic anticenter and found 15 new open clusters. To verify the reality of detected clusters, we developed a method, which is based on the assumption that probable cluster members lying along the same isochrone on the color-magnitude diagram (CMD) should also form the spatial density peak, whereas the background stars should have a flat distribution. The other benefit of this method is that it not only allows to verify a cluster candidate, but at the same time estimates the cluster's main parameters -- distance, age, and color excess. We employed this method to derive parameters of 12 new and 13 known, but poorly studied clusters, which were also detected in the field of interest, using the technique of $(J,J-H)$ and $(K_s,J-K_s)$ diagrams built with the data from 2MASS catalog. Later, $UBVI$  magnitudes of the stars in three new clusters were measured using CCD images taken at 104-cm telescope of Aryabhatta Research Institute of Observational Sciences (ARIES, India). The isochrones fitted to these data using another color-magnitude diagrams, and derived clusters parameters (Glushkova et al. 2009) show good agreement with those obtained from $(J,J-H)$ and $(K_s,J-K_s)$ 2MASS diagrams and therefore independently confirm the reality of clusters found by Koposov et al. (2008) and the correctness of the method of new cluster verification.  The aims of this study are: (1) search for overdensities in the Galactic plane in the range of the galactic latitude $|b|<24^{\\circ}$ using 2MASS data; (2) verification of some of them as a real clusters; (3) estimate of their physical parameters using proven technique we have developed earlier. ",
        "conclusions": "Using $J,H,K_s$ photometric data available in 2MASS catalog, we have performed an automated search for the star overdensities in the Galactic plane for the latitudes ranged within $|b|<24^{\\circ}$ and detected 11186 density peaks. A short list of 283 new cluster candidates has been compiled after we excluded the overdensities attributed to the background fluctuations or star concentrations coinciding with already known clusters. Employing our method described earlier (Koposov et al., 2008), we have studied these 283 candidates and verified 149 of them to be real clusters. Having involved the UKIDSS GPS survey in our investigation as an additional data source, we studied 22 more cluster candidates and found four new open clusters. Therefore, the number of new clusters found in the frame of the present paper is 153. For 130 clusters of them, we have determined their main physical parameters: distance, age, and color excess. Together with the clusters published in Koposov et al. (2008), we have found 168 new open clusters in total, and determined parameters for 142 of them. Besides that, we confirmed the existence of 14 known, but yet-unstudied clusters and determined their accurate coordinates and other parameters. Fig.~3 shows the distribution of 143 open clusters projected onto the Galactic plane. The crosses denote all clusters, except for SAI 50, whose distances are determined in this paper (129 new and 14 known clusters). We excluded cluster SAI 50 from this pattern because of large errors in its parameters. Solid circle represents the Sun. It can be clearly seen from this figure that no new clusters were discovered in the direction towards the Galactic center, because this area features maximum extinction and level of background fluctuations, which makes the detection of clusters more difficult. The major part of new clusters have been found within the radius of 2.5 kpc from the Sun, because of a relatively low limiting values of $J, H, K_s$ in the 2MASS catalog. We have not discovered any new clusters closer than 500 pc; however, we did not attempt to search for them specifically in this range of the distance, because the open cluster sample there is believed to be complete. For this reason, we used the filter allowing to detect clusters with the diameter of $10^\\prime$ and less, whereas for the close clusters this value is typically greater.  Some new rich clusters can only be seen at the IR wavelengths. For example, SAI 122 is well noticeable on the image from 2MASS data (Fig.~4) and is fully invisible on the DSS images, because the color excess in the direction toward this cluster amounts to $2^m.26$. Among the clusters discovered by us, there are few, which reveal themselves on the DSS images as a very faint objects; however, it is normally impossible to estimate their parameters. The majority of new objects are either poorly populated clusters, or rich ones, yet imperceptible against a dense background: for example, SAI 132 (see Fig.~5). In both these cases, an object cannot be detected while visually inspecting 2MASS images. Thus, the application of our automated method of searching infrared data for density peaks has provided us with the tool capable of detecting these 168 open clusters."
    },
    "0910/0910.2513_arXiv.txt": {
        "abstract": "We determine a constraint on the growth factor by measuring the damping of the baryon acoustic oscillations in the matter power spectrum using the Sloan Digital Sky Survey luminous red galaxy sample. The damping of the BAO is detected at the one sigma level. We obtain $\\sigma_8D_1(z=0.3) = 0.42^{+0.34}_{-0.28}$ at the $1\\sigma$ statistical level, where $\\sigma_8$ is the root mean square overdensity in a sphere of radius $8h^{-1}$Mpc and $D_1(z)$ is the growth factor at redshift $z$. The above result assumes that other parameters are fixed and the cosmology is taken to be a spatially flat cold dark matter universe with the cosmological constant. ",
        "introduction": "The baryon acoustic oscillations (BAO) are the sound oscillations of the primeval baryon-photon fluid prior to the recombination epoch. The BAO signature imprinted in the matter power spectrum is very useful for the study of the dark energy, hypothetically introduced to explain the accelerated expansion of the universe \\cite{Eis,Meik}, because the characteristic scale of the BAO plays a role of a standard ruler in the universe \\cite{Eis2,Seo,Matsubara}. The BAO signature in the galaxy clustering is clearly detected \\cite{Gert1,Percivalb}, and the constraints on the equation of state parameter of the dark energy are demonstrated \\cite{Percival,Gert,Okumura,Cabre}. Future large surveys for the precise measurement of the BAO are in progress or planned, providing us an important tool to explore the origin of the accelerated expansion of the universe.  Besides the measurement of the BAO, those future surveys provide us other important information. For example, the redshift-space distortions will be measured precisely at the same time. The redshift-space distortions reflect the velocity of galaxies in the direction of the line of sight. Especially, the linear redshift-space distortion due to the Kaiser effect comes from the linear velocity field of matter perturbations, whose measurement will give us a chance to test the gravity theory on the cosmological scales \\cite{Linder,Guzzo,Yamamoto,WhitePercival}. This is important because several modified gravity models are proposed to explain the accelerated expansion of the universe as an alternative to the dark energy. Thus the measurement of the redshift-space distortion, which is useful to constrain the growth rate and the growth factor of the cosmic structure, is also important.  In Ref.~\\cite{Nomura,Nomura2}, it is pointed out that the precise measurement of the BAO may also be useful to constrain the growth factor. The BAO signature observed in the galaxy power spectrum is contaminated by the nonlinear effects of the density perturbations, redshift-space distortions and the clustering bias. The effects of the nonlinearity and the redshift-space distortions cause the damping of the BAO signature compared with the case when these effects are switched off. The authors of Ref.~\\cite{Nomura} found that the damping of the BAO depends on the amplitude of the matter power spectrum and that a measurement of the BAO damping might be useful to constrain the growth factor of the cosmic structure.  In Ref.~\\cite{Nomura2}, two of the authors of the present paper reported a detection of the BAO damping using the Sloan digital sky survey (SDSS) luminous red galaxy (LRG) sample of the data release (DR) 6. Since the observed galaxy power spectrum is contaminated by errors, and is noisy, the procedure of extracting the BAO signature from a galaxy power spectrum is a subtle problem. In the previous paper~\\cite{Nomura2}, the cubic spline method was used to construct a smooth power spectrum to extract the BAO signature. However, this method is very sensitive to parameters of the cubic spline, i.e., number and interval of the nodes. Due to those reasons, it is clear that an independent cross check of the previous result is useful.  In the present paper, we adopt an independent method to construct the smooth power spectrum, which is stable compared with the cubic spline method, and re-investigate the BAO damping by confronting the model predictions with the SDSS LRG power spectrum of the DR7. We confirm the detection of the BAO damping at the one sigma level. Also we obtain a constraint on the parameter combination $D_1(z)\\sigma_8$ at the redshift $z=0.3$ from the BAO damping for the first time. Here $D_1(z)$ is the growth factor and $\\sigma_8$ is the amplitude of the root mean square overdensity in a sphere of radius $8h^{-1}$Mpc. We denote the growth rate by $f(z)=d\\ln D_1(z)/d\\ln a(z)$, where $a(z)$ is the scale factor. Throughout this paper, we use units in which the velocity of light equals 1 and the Hubble parameter $H_0=100h$km/s/Mpc with $h=0.7$. ",
        "conclusions": "In this work, we examined the BAO damping. We used the observed power spectrum of the SDSS LRG sample of DR7. The result shows that the BAO damping really exists and a constraint on $\\sigma_8D_1$ is obtained for the first time from the damping. The detection of the BAO damping is robust, and do not depend much on the specific model used to define the smooth model. The constraint on $\\sigma_8D_1$  is note very tight, but this is an independent and unique test for the growth factor. Furthermore, this method might be useful for future surveys \\cite{Nomura2}.  In the present analysis, we have not considered the effect of the clustering bias on the BAO damping. This point might be considered more carefully in the future.   \\vspace{3mm}"
    },
    "0910/0910.5796_arXiv.txt": {
        "abstract": "The majority of stars in the Galactic field and halo are part of binary or multiple systems. A significant fraction of these systems have orbital separations in excess of thousands of astronomical units, and systems wider than a parsec have been identified in the Galactic halo. These binary systems cannot have formed through the `normal' star-formation process, nor by capture processes in the Galactic field. We propose that these wide systems were formed during the dissolution phase of young star clusters. We test this hypothesis using $N$-body simulations of evolving star clusters and find wide binary fractions of $1-30$\\%, depending on initial conditions. Moreover, given that most stars form as part of a binary system, our theory predicts that a large fraction of the known wide `binaries' are, in fact, multiple systems. ",
        "introduction": "Numerous wide ($10^3$~AU$-0.1$~pc) binary systems are known to exist in the Galactic field and halo (e.g., Close et al. 1990; Chanam{\\'e} \\& Gould 2004). These binaries are usually detected as common-proper-motion pairs (e.g., Wasserman \\& Weinberg 1991; L{\\'e}pine \\& Bongiorno 2007) or through statistical methods (e.g., Bahcall \\& Soneira 1981; Garnavich 1988). Approximately 15\\% of the wide binaries in the field have a semi-major axis in the range $10^3~{\\rm AU} < a < 0.1$~pc (Duquennoy \\& Mayor 1991; see also Poveda et al. 2007). The origin of these wide binaries cannot be explained by star formation or by dynamical interactions in the Galactic field, and has long been a mystery.  The majority of stars are known to form in embedded clusters (Lada \\& Lada 2003). Wide binary systems cannot have formed through clustered star formation, simply because their semi-major axis is comparable to the size of a typical embedded cluster. Even if it were possible to form wide binaries in star clusters, they would immediately be destroyed by dynamical interactions. Although most stars are believed to have formed in star clusters, a fraction ($< 10-30$\\%) may form through diffuse (isolated) star formation. However, this fraction makes it difficult to account for the $\\sim 15$\\% wide binaries in the Galactic field. If wide binaries are formed through diffuse star formation, then the majority of stars formed in this process must form a wide binary, implying a star-formation process that is fundamentally different from that of clustered star formation.  Binary systems of any separation may also form through dynamical capture processes. To form a binary system from two interacting stars, a third star should be present to serve as energy sink. Each of the three stars should be at the same location at the same time, and should have precisely those velocities and impact angles that allow the formation of a new binary system. These conditions are extremely rare in the Galactic field, resulting only in a handful of binaries over the lifetime of the Galaxy (Goodman \\& Hut 1993). Note that binaries can form through the capture process in the cores of star clusters through this process. However, wide binary systems cannot form in these environments, simply because they have a semi-major axis comparable to (or even larger than) star cluster cores.  We propose that wide binaries are formed during the dissolution of young star clusters (Kouwenhoven et al., in prep). Young clusters generally have a short lifetime and dissolve into the field stellar population within typically 20~Myr (e.g., Mengel et al. 2005; Bastian et al. 2005; de Grijs \\& Parmentier 2007). During this dissolution process, two stars may escape the cluster in the same direction, with similar velocities. Although initially unbound, these two stars may form a new binary system after having left the gravitational potential of the cluster. In this process, the separation between the two stars is typically on the order of the size of their natal cluster at the time of dissolution, i.e., of order a parsec. ",
        "conclusions": "A significant fraction of the stars in the Galactic field and halo are part of wide ($a>10^3$~AU) binary systems. The origin of these binary systems cannot be explained by clustered star formation, nor by dynamical capture in the Galactic field and halo.  We propose that these binaries are formed during the dissolution process of young star clusters, where two initially unbound stars escape in a similar direction and form a bound wide binary (Kouwenhoven et al., in prep). We test this hypothesis using $N$-body simulations and find that we are able to reproduce wide-binary fractions of $1-30$\\% (in the semi-major-axis range $10^3~{\\rm AU} < a < 0.1$~pc), depending on the initial conditions.  Moreover, as most stars form as part of a binary system, we predict that a large fraction of the known wide `binaries' are, in fact, triple or quadruple systems, and the ratio of wide binary, triple and quadruple systems therefore provides an indication of the primordial binary fraction. \\newpage"
    },
    "0910/0910.0450_arXiv.txt": {
        "abstract": "{A second large programme (LP) for the physical studies of TNOs and Centaurs, started at ESO Cerro Paranal on October 2006 to obtain high-quality data, has recently been concluded.  In this paper we present the spectra of these pristine bodies obtained in the visible range during the last two semesters (November 2007--November 2008) of the LP.}{ We investigate the spectral behaviour of the TNOs and Centaurs observed, and we analyse the spectral slopes distribution of the full data set coming from this LP and from the literature.} {Spectroscopic observations in the visible range were carried out at the UT1 (Antu) telescope using the instrument FORS2. We computed the spectral slope for each observed object, and searched for possible weak absorption features. A statistical analysis was performed on a total sample of 73 TNOs and Centaurs to look for possible correlations between dynamical classes, orbital parameters, and spectral gradient. } {We obtained new spectra for 28 bodies (10 Centaurs, 6 classical, 5 resonant, 5 scattered disk, and 2 detached objects), 15 of which were observed for the first time. All the new presented spectra are featureless, including 2003 AZ84, for which a faint and broad absorption band possibly attributed to hydrated silicates on its surface has been reported (Fornasier et al. 2004a, and Alvarez-Candal et al. 2008). The data confirm a wide variety of spectral behaviours, with neutral--grey to very red gradients. An analysis of the spectral slopes available from this LP and in the literature for a total sample of 73 Centaurs and TNOs shows that there is a lack of very red objects in the classical population. We present the results of the statistical analysis of the spectral slope distribution versus orbital parameters. In particular, we confirm a strong anticorrelation between spectral slope and orbital inclination for the classical population. Nevertheless, we do not observe a change in the slope distribution at $i \\sim 5^{o}$, the boundary between the dynamically hot and cold populations, but we find that objects with $ i < 12^{o}$ show no correlation between spectral slope and inclination, as already noticed by Peixinho et al. (2008) on the colour --inclination relation for classical TNOs. A strong correlation is also found between the spectral slope and orbital eccentricity for resonant TNOs, with objects having higher spectral slope values with increasing eccentricity.}{} ",
        "introduction": "The icy bodies in orbit beyond Neptune (trans-Neptunians or TNOs) and Centaurs represent the most primitive population of all the Solar System objects. They are fossil remnants of the formation of our planetary system and the investigation of their surface properties is essential for understanding the formation and the evolution of the Solar System. Since 1992, more than 1300 TNOs and Centaurs have been detected. The TNO population is dynamically classified into several categories (Gladman et al. 2008): classical objects have orbits with low eccentricities and semi-major axes between about 42 and 48 AU; resonant objects are trapped in resonances with Neptune, the majority of them located in or near the 3:2 mean motion resonance; scattered objects have high-eccentricity, high-inclination orbits and a perihelion distance near q=35 AU; extended scattered disk (or detached) objects (SDOs) are located at distances so great that they cannot have been emplaced by gravitational interactions with Neptune. In addition, the Centaurs are closest to the Sun and have unstable orbits between those of Jupiter and Neptune. They seem to originate from the Kuiper Belt and should have been injected into their present orbits by perturbations from giant planets or mutual collisions.  The investigation of the surface composition of these icy bodies can provide essential information about the conditions in the early Solar System environment at large distances from the Sun. \\\\ Aiming at obtaining high--quality data of these populations, a large programme (LP) for the observation of Centaurs and TNOs was started using the facilities of the ESO / Very Large Telescope site at Cerro Paranal in Chile (PI: M.A. Barucci). More than 500 hours were allocated to observe these objects in visible and near-infrared photometry--spectroscopy, and V--R polarimetry from November 2006 until December 2008.  In this paper we report the results of the visible spectroscopy of TNOs and Centaurs performed with the instrument FORS2 at the VLT unit 1 {\\it Antu} during the last two semesters of the LP (November 2007-November 2008). New visible spectra of 28 objects with $ 18.7 < V < 21.8$, were acquired, 15 of which were observed for the first time. A statistical analysis of the full data set of spectral slopes coming from this LP and from the literature is presented on a total sample of 73 TNOs and Centaurs.  \\begin{table*}[t] \\caption{Observational conditions of the TNOs and Centaurs spectroscopically investigated. } {\\footnotesize \\label{tab_obs} \\begin{center} \\begin{tabular}{|l|l|l|c|c|c|c|c|l|c|} \\hline {\\bf  Object} & {\\bf Date} & {\\bf UT} &  {\\boldmath{$\\Delta$}} & {{\\boldmath $r$}} & {\\boldmath $\\alpha$} &{\\bf T\\boldmath{$_{exp}$}} & {\\bf airm.} & {\\bf Solar An. (airm.)} & {\\bf Slope} \\\\ & & &   {\\bf [AU]} & {\\bf [AU]} &  {\\boldmath [$^{o}$]} & {\\bf [s]} &  & & {\\bf [\\%/(\\boldmath{$10^{3}$ \\AA)}]} \\\\ \\hline {\\em CENTAURS} & & & & & & & & & \\\\ 5145 Pholus        & 2008 Apr. 12  & 05:57:44 & 21.214 & 21.864 & 2.0 & 2600 & 1.21 & Ld102$-$1081 (1.15) &  48.6$\\pm$0.7 \\\\ 10199 Chariklo     & 2008 Feb. 04  & 07:57:47 & 13.301 & 13.395 & 4.2 & 900 & 1.10 & Ld102$-$1081 (1.13)  &  10.4$\\pm$0.6 \\\\ 52872 Okyrhoe      & 2008 Feb. 04  & 06:33:30 &  4.879 &  5.800 & 3.8 & 1800 & 1.14 & Ld102$-$1081 (1.13)  &  13.4$\\pm$0.6\\\\ 55576 Amycus       & 2008 Apr. 12  & 03:43:14 & 15.206 & 16.056 & 2.0 & 2400 & 1.16 & Ld102$-$1081 (1.15)  &  37.1$\\pm$0.9 \\\\ {\\bf 120061 (2003 CO1)}  & 2008 Apr. 12  & 04:48:44 & 10.180 & 11.122 & 1.8 & 2400 & 1.13 & Ld102$-$1081 (1.15)  &  9.4$\\pm$0.6  \\\\ {\\bf 2002 KY14}          & 2008 Sept. 21 & 01:11:40 &  7.802 &  8.649 & 3.8 & 1200 & 1.25 & Ld112$-$1233(1.12)   & 36.9$\\pm$0.7  \\\\ {\\bf 2007 UM126}         & 2008 Sept. 21 & 04:46:01 & 10.202 & 11.163 & 1.6 & 2400 & 1.10 & Ld112$-$1233(1.12)   &  8.8$\\pm$0.7 \\\\ {\\bf 2007 VH305}         & 2008 Nov. 22  & 02:38:07 &  7.853 &  8.638 & 4.2 & 2400 & 1.18 & Ld98$-$978 (1.20)   & 12.1$\\pm$0.7 \\\\ {\\bf 2008 FC76}          & 2008 Sept. 21 & 00:21:29 & 10.976 & 11.688 & 3.6 & 1800 & 1.29 & Ld112$-$1233(1.12)   & 36.0$\\pm$0.7 \\\\ {\\bf 2008 SJ236}         & 2008 Nov. 22  & 00:53:58 &  5.522 &  6.364 & 5.0 & 4000 & 1.28 & Ld93$-$101 (1.30)   & 20.9$\\pm$0.8 \\\\ \\hline {\\em CLASSICALS} & & & & & & & & & \\\\ {\\bf 20000 Varuna}       & 2007 Dec. 04  & 06:41:41 & 42.606 & 43.392 & 0.8 & 2400 & 1.57 & Hyades64 (1.41)      &  24.9$\\pm$0.6     \\\\ 120178 (2003 OP32) & 2008 Sept. 21 & 02:26:48 & 40.545 & 41.365 & 0.8 & 2400 & 1.14 & Ld112$-$1233(1.12)   &  0.9$\\pm$0.7  \\\\ {\\bf 120348 (2004 TY364) }& 2008 Nov. 22  & 04:53:01 & 38.839 & 39.591 & 0.9 & 2400 & 1.27 & Ld93$-$101 (1.30)   & 22.9$\\pm$0.7 \\\\ {\\bf 144897 (2004 UX10)} & 2007 Dec. 04  & 01:10:32 & 38.144 & 38.836 & 1.0 & 2400 & 1.13 & HD1368 (1.12)        &  20.7$\\pm$0.8  \\\\ {\\bf 144897 (2004 UX10)} & 2007 Dec. 05  & 02:41:59 & 38.158 & 38.836 & 1.1 & 2400 & 1.19 & Ld93101 (1.24)       &  19.2$\\pm$0.9  \\\\ 145453 (2005 RR43) & 2007 Dec. 04  & 02:59:10 & 37.623 & 38.511 & 0.6 & 2400 & 1.12 & HD1368 (1.12)        &   1.6$\\pm$0.6   \\\\ {\\bf 174567 (2003 MW12)} & 2008 Apr. 12  & 07:17:41 & 47.318 & 47.968 & 0.9 & 3700 & 1.09 & Ld102$-$1081 (1.10)  &  19.2$\\pm$0.6 \\\\ \\hline {\\em RESONANTS} & & & & & & & & & \\\\ {\\bf 42301 (2001 UR163)}  & 2007 Dec. 05  & 01:10:23 & 49.625 & 50.308 & 0.8 & 2400 & 1.23 & Ld93$-$101 (1.24)    &  50.9$\\pm$0.7     \\\\ 55638 (2002 VE95)  & 2007 Dec. 05  & 04:13:26 & 27.297 & 28.248 & 0.5 & 2400 & 1.21 & Ld93$-$101 (1.24)    &  40.0$\\pm$0.7      \\\\ 90482 Orcus        & 2008 Feb. 04  & 05:25:24 & 46.901 & 47.807 & 0.5 & 2400 & 1.07 & Ld102$-$1081 (1.11)  &   1.6$\\pm$0.6  \\\\ 208996 (2003 AZ84) & 2008 Nov. 22  & 06:29:23 & 44.883 & 45.458 & 1.0 & 2800 & 1.30 & Ld93$-$101 (1.30)   & 3.6$\\pm$0.6 \\\\ {\\bf 2003 UZ413}         & 2007 Dec. 04  & 04:42:24 & 41.171 & 42.004 & 0.7 & 2400 & 1.36 & Hyades64 (1.41)      &   6.2$\\pm$0.6  \\\\ {\\bf 2003 UZ413}         & 2008 Sept. 21 & 05:46:16 & 41.469 & 42.163 & 1.0 & 2400 & 1.25 & Ld112$-$1233(1.12)   &  5.0$\\pm$0.8  \\\\ {\\bf 2003 UZ413}         & 2008 Nov. 22  & 03:49:31 & 41.275 & 42.197 & 0.5 & 2400 & 1.15 & Ld98$-$978 (1.17)   & 4.9$\\pm$0.7 \\\\ \\hline {\\em SDOs} & & & & & & & & & \\\\ 42355 Typhon       & 2008 Apr. 12  & 01:07:02 & 16.892 & 17.650 & 2.2 & 1200 & 1.22 & Ld102$-$1081 (1.15)  &  12.1$\\pm$0.8 \\\\ 44594 (1999 OX3)   & 2008 Sept. 21 & 03:22:39 & 22.023 & 22.889 & 1.3 & 2400 & 1.07 & Ld110$-$361 (1.10)   & 36.2$\\pm$0.8  \\\\ {\\bf 73480  (2002 PN34)} & 2007 Nov. 10  & 00:51:51 & 14.893 & 15.344 & 3.3 & 2800 & 1.22 & LD115$-$871 (1.13)   &  15.8$\\pm$0.7 \\\\ {\\bf 145451 (2005 RM43)} & 2007 Dec. 04  & 04:04:16 & 34.298 & 35.195 & 0.7 & 2400 & 1.15 & HD1368 (1.12)        &   2.2$\\pm$0.7    \\\\ {\\bf 2007 UK126}         & 2008 Sept. 21 & 08:18:31 & 45.136 & 45.619 & 1.1 & 2400 & 1.07 & Ld112$-$1233(1.12)   & 19.6$\\pm$0.7 \\\\ \\hline {\\em DETACHED OBJECTS} & & & & & & & & & \\\\ 90377 Sedna        & 2008 Sept. 21 & 06:53:28 & 87.419 & 88.015 & 0.5 & 3600 & 1.20 & Ld112$-$1233(1.12)   & 40.2$\\pm$0.9 \\\\ 120132 (2003 FY128)& 2008 Apr. 12  & 02:32:57 & 37.477 & 38.454 & 0.3 & 2800 & 1.06 & Ld102$-$1081 (1.10)  &  26.7$\\pm$1.0  \\\\ \\hline \\end{tabular} \\end{center} } \\begin{list}{}{} \\item Observational conditions: for each object we report the observational date and universal time (UT of the beginning of the exposure), the geocentric distance ($\\Delta$), the heliocentric distance ($r$), the phase angle ($\\alpha$), the total exposure time, the airmass (mean of the airmass value at the beginning and at the end of observation), the observed solar analogue stars with their airmass used to remove the solar contribution, and the spectral slope value computed in the 0.5--0.8 $\\mu$m range. The TNOs and Centaurs observed in spectroscopy for the first time are in bold. \\end{list} \\end{table*} ",
        "conclusions": "In this work we have presented new visible spectra for 18 TNOs and 10 Centaurs obtained in the framework of a LP devoted to the investigation of these pristine bodies. The main results obtained may be summarised as follows. \\begin{itemize} \\item Our data confirm the diversity in the surface composition of the Centaur and TNO populations, with a spectral slope gradient ranging from $\\sim$1 to 51 \\%/10$^3$ \\AA. \\item Absorption features: all the new spectra are featureless, excluding 10199 Chariklo and 42355 Typhon, which show faint broadband absorptions centred at 0.62-0.65 $\\mu$m, tentatively attributed to phyllosilicates. These two objects are fully discussed in two separate papers (Guilbert et al. 2009, Alvarez-Candal et al. 2009) \\item The new spectrum obtained of the plutino 2003 AZ84 is featureless, while a faint broadband feature possibly attributed to hydrated silicates was previously reported by Fornasier et al. (2004) and Alvarez-Candal (2008), so we cannot confirm an aqueous alteration action on its surface, but we cannot exclude that 2003 AZ84 simply has a heterogeneous surface composition. \\item Spectral heterogeneity: comparing the spectral slope values of the objects obtained here with those available in the literature, we found important differences for a few objects, in particular for 10199 Chariklo, 20000 Varuna, 44594 (1999 OX3), and 120132 (2003 FY128). This variability on the spectral behaviour may stem from surface heterogeneities. We observed the fast-rotating (period of 4.13$\\pm$0.05 hours, Perna et al. 2009b) plutino 2003 UZ413 during 3 different runs. We found that the slope of the spectrum acquired on the second maximum of the lightcurve is higher than the two slopes obtained near the minima, even if the differences are not greater than the uncertainties. Further observations covering the whole rotational period of 2003 UZ413 are needed to confirm if the spectral variation on its surface is real and if this object may display surface variability. \\item  Spectral slope distribution: we combined the new data presented here with the visible spectra published in the literature (coming mainly from the 2 LPs devoted to TNOs and Centaurs carried out at ESO). The total sample contains 20 Centaurs and 53 TNOs. After analysing the spectral slope distribution of Centaurs, classicals, resonants, and SDOs--detached bodies, we found a lack of very red objects in the classical population, as all have spectral slope values $<$ 35 \\%/10$^3$ \\AA, with the lowest mean value ($\\sim$ 14 \\%/10$^3$ \\AA) in the dynamical classes investigated. The Centaur population is dominated (70\\%) by bluish to grey or moderately red objects but also comprises very red objects, with spectral slopes greater than 40\\%/10$^3$ \\AA. Nethertheless, very red objects are a minority in the Centaur and TNO populations (8\\%), while most of them (74\\%) have spectral slope values $<$ 25\\%/10$^3$ \\AA. \\item Statistical analysis: running a Spearman rank order analysis, we confirm a strong anticorrelation between the spectral slope and inclination for the classical TNOs, of a weak anticorrelation between $S$ and $e$, and of a weak correlation between $S$ and the perihelion distance. Our sample mainly contains objects in the hot population ($i> 5^{o}$), which thus show neutral/blue spectral behaviour for increasing values of $i$ and $e$. Nevertheless, we do not observe a change in the slope distribution between the hot and cold populations, but we do find that objects with $ i < 12^{o}$ show no correlation between spectral slope and inclination. This result independently confirms the finding by Peixinho et al. (2008) on the colour--inclination relation for classical TNOs. We found a strong correlation between $S$ and $e$ for resonant bodies, a similar but weaker one for Centaurs, while for SDOs we confirm the anticorrelation between $S$ and $e$ found by Santos-Sanz et al. (2009), even though it must be noted that our sample is very limited. \\end{itemize}"
    },
    "0910/0910.3948_arXiv.txt": {
        "abstract": "The resistive decay of chains of three interlocked magnetic flux rings is considered. Depending on the relative orientation of the magnetic field in the three rings, the late-time decay can be either fast or slow. Thus, the qualitative degree of tangledness is less important than the actual value of the linking number or, equivalently, the net magnetic helicity. Our results do not suggest that invariants of higher order than that of the magnetic helicity need to be considered to characterize the decay of the field. ",
        "introduction": "\\label{Introduction}  Magnetic helicity plays an important role in plasma physics \\cite{Tay74,JC84,BF84}, solar physics \\cite{Low96,RK94,RK96}, cosmology \\cite{BEO96,FC00,CHB05}, and dynamo theory \\cite{PFL76,BS05}. This is connected with the fact that magnetic helicity is a conserved quantity in ideal magnetohydrodynamics  \\cite{Wol58}. The conservation law of magnetic helicity is ultimately responsible for inverse cascade behavior that can be relevant for spreading primordial magnetic field over large length scales. It is also likely the reason why the magnetic fields of many astrophysical bodies have length scales that are larger than those of the turbulent motions responsible for driving these fields. In the presence of finite magnetic diffusivity, the magnetic helicity can only change on a resistive time scale. Of course, astrophysical bodies are open, so magnetic helicity can change by magnetic helicity fluxes out of or into the domain of interest. However, such cases will not be considered in the present paper.  In a closed or periodic domain without external energy supply, the decay of a magnetic field depends critically on the value of the magnetic helicity. This is best seen by considering spectra of magnetic energy and magnetic helicity. The magnetic energy spectrum $M(k)$ is normalized such that \\EQ \\int M(k)\\,d k=\\bra{\\BB^2}/2\\mu_0, \\EN where $\\BB$ is the magnetic field, $\\mu_0$ is the magnetic permeability, and $k$ is the wave number (ranging from 0 to $\\infty$). The magnetic helicity spectrum $H(k)$ is normalized such that \\EQ \\int H(k)\\,d k=\\bra{\\AAA\\cdot\\BB}, \\EN where $\\AAA$ is the magnetic vector potential with $\\BB=\\nab\\times\\AAA$. In a closed or periodic domain, $H(k)$ is gauge-invariant, i.e.\\ it does not change after adding a gradient term to $\\AAA$. For finite magnetic helicity, the magnetic energy spectrum is bound from below \\cite{Wol58} such that \\EQ M(k)\\ge k|H(k)|/2\\mu_0. \\label{bound} \\EN This relation is also known as the realizability condition \\cite{Mof69}. Thus, the decay of a magnetic field is subject to a corresponding decay of its associated magnetic helicity. Given that in a closed or periodic domain the magnetic helicity changes only on resistive time scales \\cite{Ber84}, the decay of magnetic energy is slowed down correspondingly. More detailed statements can be made about the decay of turbulent magnetic fields, where the energy decays in a power-law fashion proportional to $t^{-\\sigma}$. In the absence of magnetic helicity, $\\bra{\\AAA\\cdot\\BB}=0$, we have a relatively rapid decay with $\\sigma\\approx1.3$ \\cite{MacLow}, while with $\\bra{\\AAA\\cdot\\BB}\\neq0$, the decay is slower with $\\sigma$ between 1/2 \\cite{CHB05} and  2/3 \\cite{BM99}.  The fact that the decay is slowed down in the helical case is easily explained in terms of the topological interpretation of magnetic helicity. It is well known that the magnetic helicity can be expressed in terms of the linking number $n$ of discrete magnetic flux ropes via \\cite{Mof69} \\EQ \\int\\AAA\\cdot\\BB\\,\\dd V=2n\\Phi_1\\Phi_2, \\EN where \\EQ \\Phi_i=\\int_{S_i}\\BB\\cdot\\dd\\SSS\\quad\\mbox{(for $i=1$ and 2)} \\label{FluxDef} \\EN are the magnetic fluxes of the two ropes with cross-sectional areas $S_1$ and $S_2$. The slowing down of the decay is then plausibly explained by the fact that a decay of magnetic energy is connected with a decay of magnetic helicity via the realizability condition \\eqref{bound}. Thus, a decay of magnetic helicity can be achieved either by a decay of the magnetic flux or by magnetic reconnection. Magnetic flux can decay through annihilation with oppositely oriented flux. Reconnection on the other hand reflects a change in the topological connectivity, as demonstrated in detail in Ref.~\\cite[p.28]{PriestRecon2000}.  The situation becomes more interesting when we consider a flux configuration that is interlocked, but with zero linking number. This can be realized quite easily by considering a configuration of two interlocked flux rings where a third flux ring is connected with one of the other two rings such that the total linking number becomes either 0 or 2, depending on the relative orientation of the additional ring, as is illustrated in \\Fig{flux_configuration}. Topologically, the configuration with linking numbers of 0 and 2 are the same except that the orientation of the field lines in the upper ring is reversed. Nevertheless, the simple topological interpretation becomes problematic in the case of zero linking number, because then also the magnetic helicity is zero, so the bound of $M$ from below disappears, and $M$ can now in principle freely decay to zero. One might expect that the topology should then still be preserved, and that the linking number as defined above, which is a quadratic invariant, should be replaced with a higher order invariant \\cite{RA94,HM02,Kom09}. It is also possible that in a topologically interlocked configuration with zero linking number the magnetic helicity spectrum $H(k)$ is still finite and that bound \\eq{bound} may still be meaningful. In order to address these questions we perform numerical simulations of the resistive magnetohydrodynamic equations using simple interlocked flux configurations as initial conditions. We also perform a control run with a non-interlocked configuration and zero helicity in order to compare the magnetic energy decay with the interlocked case.  \\begin{figure}[t!]\\begin{center} \\includegraphics[width=.25\\columnwidth]{H0_flux_configuration} \\qquad \\includegraphics[width=.25\\columnwidth]{H4_flux_configuration} \\qquad \\includegraphics[width=.3\\columnwidth]{H0_flux_configuration_no_links} \\end{center}\\caption[]{ Visualization of the triple ring configuration at the initial time. Arrows indicate the direction of the field lines in the rings, corresponding to a configuration with $n=0$ (left) and $n=2$ (center). On the right the non-interlocked configuration with $n=0$ is shown. }\\label{flux_configuration}\\end{figure}  Magnetic helicity evolution is independent of the equation of state and applies hence to both compressible and incompressible cases. In agreement with earlier work \\cite{KB99} we assume an isothermal gas, where pressure is proportional to density and the sound speed is constant. However, in all cases the bulk motions stay subsonic, so for all practical purposes our calculations can be considered nearly incompressible, which would be an alternative assumption that is commonly made \\cite{GM00}. ",
        "conclusions": "The present work has shown that the rate of magnetic energy dissipation is strongly constrained by the presence of magnetic helicity and not by the qualitative degree of knottedness. In our example of three interlocked flux rings we considered two flux chains, where the topology is the same except that the relative orientation of the magnetic field is reversed in one case. This means that the linking number switches from 2 to 0, just depending on the sign of the field in one of the rings. The resulting decay rates are dramatically different in the two cases, and the decay is strongly constrained in the case with finite magnetic helicity.  The present investigations reinforce the importance of considering magnetic helicity in studies of reconnection. Reconnection is a subject that was originally considered in two-dimensional studies of X-point reconnection \\cite{Par57,Bis86}. Three-dimensional reconnection was mainly considered in the last 20 years. An important aspect is the production of current sheets in the course of field line braiding \\cite{Ber93}. Such current sheets are an important contributor to coronal heating \\cite{GN96}. The crucial role of magnetic helicity has also been recognized in several papers \\cite{Hu97,Liu07}. However, it remained unclear whether the decay of interlocked flux configurations with zero helicity might be affected by the degree of tangledness. Our present work suggests that a significant amount of dissipation should only be expected from tangled magnetic fields that have zero or small magnetic helicity, while tangled regions with finite magnetic helicity should survive longer and are expected to dissipate less efficiently."
    },
    "0910/0910.3665_arXiv.txt": {
        "abstract": "We investigate the effects of the Cosmic Microwave Background (CMB) radiation field on the collapse of prestellar clouds. Using a semi-analytic model to follow the thermal evolution of clouds with varying initial metallicities and dust contents at different redshifts, we study self-consistently the response of the mean Jeans mass at cloud fragmentation to metal line-cooling, dust-cooling and the CMB.  In the absence of dust grains, at redshifts $z \\le 10$ moderate characteristic masses (of 10s of $M_{\\odot}$) are formed when the metallicity is $10^{-4} Z_{\\odot} \\le Z \\le 10^{-2.5} Z_{\\odot}$; at higher metallicities, the CMB inhibits fragmentation and only very large masses (of $\\sim$ 100s of $M_{\\odot}$) are formed. These effects become even more dramatic at $z > 10$ and the fragmentation mass scales are always $\\ge$ 100s of $M_{\\odot}$, independent of the initial metallicity.  When dust grains are present, sub-solar mass fragments are formed at any redshift for metallicities $Z \\ge 10^{-6} Z_{\\odot}$ because dust-cooling remains relatively insensitive to the presence of the CMB. When $Z > 10^{-3} Z_{\\odot}$, heating of dust grains by the CMB at $z \\ge 5$ favors the formation of larger masses, which become super-solar when $Z \\ge 10^{-2} Z_{\\odot}$ and $z \\ge 10$. Finally, we discuss the implications of our result for the interpretation of the observed abundance patterns of very metal-poor stars in the galactic halo. ",
        "introduction": "The origin of the stellar Initial Mass Function (IMF) and its variation with cosmic time or with diverse environmental conditions still lack a complete physical interpretation. Observationally, the present-day stellar IMF appears to have an almost universal profile, characterized by a power-law at large masses and flattening below a characteristic mass of $M_{\\rm ch} \\approx 1 M_\\odot$ with a plateau which extends down to the brown dwarf regime (see Chabrier 2003 and references therein). This mass scale, which represents the \"typical\" outcome of the star formation process, appears to be remarkably uniform in diverse environments where the initial conditions for star formation can vary considerably (see Elmegreen, Klessen \\& Wilson 2008).  Among the many proposed explanations, the origin of the characteristic stellar mass and the broad peak of the IMF are best attributed to gravitational fragmentation, which sets the mean Jeans mass at cloud fragmentation (see Bonnell, Larson \\& Zinnecker 2007 for a comprehensive review), \\begin{equation} M_{\\rm frag} = \\rho \\lambda_J^3 \\sim 700 M_{\\odot} \\left(\\frac{T_{\\rm frag}}{200~\\rm{K}}\\right)^{3/2} \\left(\\frac{n_{\\rm frag}}{10^4 \\rm{cm}^{-3}}\\right)^{-1/2} \\label{eq:mfrag} \\end{equation} where the Jeans length is defined as, \\[ \\lambda_J = \\left(\\frac{\\pi k T}{\\mu m_H G \\rho}\\right)^{1/2}. \\] \\noindent The process of gravitational fragmentation can be investigated using thermal physics within the framework of the classical Jeans criterion: fragmentation occurs at the end of an efficient cooling phase, when the temperature decreases for increasing density and the effective adiabatic index switches from $\\gamma < 1$ to $\\gamma > 1$ (the pressure $p \\propto \\rho^{\\gamma}$, where $\\rho$ is the density); the typical fragmentation mass-scale is given by the thermal Jeans mass at that epoch (Larson 2005; Schneider et al. 2002). The appearance of a characteristic mass at the inflection point in the equation of state has been also demonstrated numerically in 3D simulations by Jappsen et al. (2005), supporting the idea that the distribution of stellar masses depends mainly on the thermodynamical state of the star-forming gas.  In the physical conditions which apply to star forming regions in the present-day Universe, the application of the above criterion leads to a characteristic fragmentation mass of $M_{\\rm frag} \\sim 2 M_\\odot$. About a half of this gas is expected to be ejected by a protostellar outflow without accreting (e.g. Matzner \\& McKee 2000).  If we apply the same criterion to primordial star forming regions we find a fragmentation mass of order $M_{\\rm frag} \\sim 700 M_\\odot$, which corresponds to the Jeans mass at the temperature and density where cooling by molecular hydrogen becomes less efficient due to the NLTE-LTE transition in the level populations of the molecule. These clumps are the immediate progenitors of the stars, in agreement with the results found by 3D numerical simulations (Abel, Bryan \\& Norman 2002; Bromm, Coppi \\& Larson 2002; Yoshida, Omukai \\& Hernquist 2008).  Thus, between the primordial and the present-day conditions there is a clear transition in fragmentation scales which is believed to be mainly driven by the progressive enrichment of star forming clouds with metals (Bromm et al. 1999; Omukai 2000; Schneider et al. 2002; Bromm \\& Loeb 2003; Santoro \\& Shull 2006; Smith, Sigurdsson \\& Abel 2008) and dust grains (Schneider et al. 2003; Omukai et al. 2005; Tsuribe \\& Omukai 2006; Clark, Glover \\& Klessen 2008).  Note, however, that the presence of gas-phase metals can not by itself lead to solar or sub-solar characteristic masses. In fact, metal line-cooling is efficient in the initial evolution of collapsing gas clouds, for densities in the range $10^4 - 10^6~$cm$^{-3}$, and induces fragmentation into relatively large clumps, with characteristic fragment masses of a few 10s of $M_\\odot$ up to initial gas metallicities $Z \\lsim 10^{-2} Z_{\\odot}$ (Schneider et al. 2006). Conversely, the formation of solar or sub-solar mass fragments might operate already at metallicities $Z_{\\rm cr} = 10^{-6} Z_{\\odot}$ in the presence of dust grains, which provide an efficient source of cooling at densities $\\gsim 10^{13}$cm$^{-3}$. Since dust can be promptly synthesized in the ejecta of primordial supernovae (Todini \\& Ferrara 2001; Nozawa et al. 2003; Schneider, Ferrara \\& Salvaterra 2004; Bianchi \\& Schneider 2007), it is likely that its contribution to thermal and fragmentation history of collapsing clouds at high redshifts had been relevant.  An additional redshift-dependent effect is induced by the Cosmic Microwave Background (CMB) which sets an effective temperature floor increasing with redshift. Clarke \\& Bromm (2003) have proposed a simple formula for the characteristic stellar mass which is identified by the Jeans mass fixed jointly by the gas temperature and the pressure set by the weight of the overlying baryons as the gas collapses in the parent dark matter halo. In this simple model, they do not follow the gas thermal evolution and the temperature is taken to be the maximum value between the CMB floor and either $200$~K for H$_2$-cooling or $10$~K for CO-cooling. In the latter case, the characteristic stellar mass is set by the CMB temperature at all redshifts $z > 2.7$ and it is predicted to be $> 10 M_{\\odot}$ at $z > 10$.  The CMB influences the thermal evolution of collapsing gas clouds in two possible ways: (i) it affects the atomic and molecular line level populations (ii) it heats the dust grains. In Omukai et al. (2005), we have estimated the impact of the CMB temperature floor at $z = 20$ on the evolution of prestellar clouds with varying initial metallicity. It was shown that the CMB affects the cloud evolution at low temperatures, which are reached for relatively large values of the initial metallicity, $Z > 10^{-3} Z_\\odot$.  Smith et al. (2009) have investigated the effects of the CMB on the collapse of prestellar clouds with varying initial metallicities by means of numerical simulations. They confirm that due to metal line-cooling, the gas becomes thermally coupled to the CMB for $Z > 10^{-2.5} Z_\\odot$ early in the evolution of the clouds when the density is still low. Note that a similar effect has been recently found by Jappsen et al. (2009) using 3D simulations starting from different initial conditions, namely protogalaxies forming within a previously ionized HII region that has not yet had time to cool and recombine. Even in the latter models, the gas cools rapidly to the CMB floor and no fragmentation occurs during the collapse for metallicities as high as $0.1 Z_{\\odot}$.  Smith et al. (2009) conclude that the CMB inhibits fragmentation at large metallicities and that there were three distinct modes of star formation at high redshift: a \u2018primordial\u2019 mode, producing very massive stars at very low metallicities ($Z < 10^{-3.75} Z_\\odot$); a CMB-regulated mode, producing moderate mass (10s of $M_\\odot$) stars at high metallicites ($Z > 10^{-2.5} Z_\\odot$ at redshift $z \\sim 15-20$); and a low-mass (a few $M_{\\odot}$) mode existing between those two metallicities. As the Universe ages and the CMB temperature decreases, the range of the low-mass mode extends to higher metallicities, eventually becoming the only mode of star formation.  In their simulations, however, Smith et al. (2009) consider the thermal evolution in the presence of molecular coolants (H$_2$ and HD) and gas-phase metals only. They do not include H$_2$ formation heating, which affects the thermal evolution at moderate densities $\\geq 10^8$cm$^{-3}$ in the primordial gas and at lower densities in the metal-enriched gas. Moreover, they do not take into consideration the possible presence of dust grains which significantly affect the gas thermal evolution even at very low metallicities via cooling by thermal emission and efficient H$_2$ formation on their surfaces (Schneider et al. 2003; Omukai et al. 2005; Schneider et al. 2006).  From an observational point of view, the number and properties of very metal-poor (VMP) stars in the Galactic halo with [Fe/H] $<-2$ seem to require a higher characteristic stellar mass in the past, with a redshift-modulation consistent with that expected by the CMB (Hernandez \\& Ferrara 2001). Recent analyses of the surface elemental composition of very metal-poor halo stars (Cayrel et al. 2004) show that the observed abundance pattern appears consistent with the predicted yields of Pop III stars with characteristic mass of $\\approx 10 M_{\\odot}$ (Heger \\& Woosley 2008). Similar conclusions have been derived analysing the fraction of carbon-enhanced stars (CEMP, Tumlinson 2007; Komiya et al 2007). CEMP stars are a subset of metal-poor stars that show enhanced carbon-to-iron abundances ([C/Fe] $> 1$) (Beers \\& Christlieb 2005). These stars are thought to be metal-poor low-mass stars ($< 0.8 M_\\odot$) that have acquired C enhancements at their convective surfaces by capturing the C-rich ejecta of an AGB companion ($1.5 M_\\odot <  M_\\star < 8 M_\\odot$). Because the binary system that produces a CEMP star requires both a low-mass star and an intermediate mass star, the incidence and chemical abundance signatures of CEMP stars reflects the underlying IMF in the range $1 M_\\odot < M_\\star < 8 M_\\odot$ (Lucatello et al. 2005; Tumlinson 2007; Komiya et al. 2007).  On the basis of these considerations, Tumlinson (2007) suggests that the CMB is shaping the characteristic mass of the IMF and proposes the following parametrization, \\begin{equation} M_{\\rm ch}/M_\\odot = 1.06 M_\\odot \\left(\\frac{{\\rm max}[(1+z)2.73~{\\rm K}, 8~{\\rm K}]}{10~{\\rm K}}\\right)^{1.7}, \\end{equation} \\noindent which implies that $M_{\\rm ch} = 0.72, 6.87$ and $20.63 M_\\odot$ at $z = 0, 10$ and $20$.  The question is whether this observationally motivated CMB-regulated IMF can be reconciled with theoretical models for the thermal evolution and fragmentation of star-forming clouds at very low metallicities.  Here we intend to revisit the role of the CMB in the evolution of protostellar clouds investigating in a self-consistent way the response of the fragmentation mass scale to metal line-cooling, dust cooling, and the CMB.  The paper is organized as follows: in section \\ref{sec:model} we present the model adopted to follow the thermal and chemical evolution of collapsing protostellar clouds, in section \\ref{sec:results} we illustrate the main results of the analysis, and in section \\ref{sec:conclusions} we discuss their implications. ",
        "conclusions": "\\label{sec:conclusions}  In this paper we have investigated the dependence of thermal and fragmentation properties of star-forming clouds on the environmental conditions, focussing on the interplay between metals, dust grains and the CMB at different redshifts.  Our results, based on an extended grid of models with varying initial metallicity and redshift, can be summarized as follows:  \\begin{enumerate}  \\item In the absence of dust grains, when fragmentation occurs at the end of the line-cooling phase, moderate characteristic masses (of 10s of $M_{\\odot}$) are formed at redshifts $z \\le 10$ when the metallicity is $10^{-4} Z_{\\odot} \\le Z \\le 10^{-2.5} Z_{\\odot}$; at lower metallicities, the clouds follow the primordial fragmentation mode and at higher metallicity, $Z > 10^{-2.5} Z_{\\odot}$, the CMB inhibits fragmentation and only very large masses $\\sim$ 100s of $M_{\\odot}$, are formed. These effects become even more dramatic at $z > 10$, when the fragmentation mass scales are always $\\ge$ 100s of $M_{\\odot}$, independent of the initial metallicity.  \\item When dust grains are present, at $z=0$ two regimes are present: the line-induced fragmentation mode, which leads to fragment masses $> 50 M_{\\odot}$ at any metallicity, and the dust-induced fragmentation mode, which occurs at higher densities and leads to sub-solar mass fragments when $Z \\ge 10^{-6} Z_{\\odot}$. The effects of the CMB at $z > 0$ are important in both physical regimes: at very low metallicities, $Z < 10^{-4} Z_{\\odot}$, line-cooling becomes less efficient because of the suppression of H$_2$ formation on dust grains when the grain temperature becomes $\\ge 75$~K, at $z > 20$. At higher metallicities, metal line-cooling causes the gas to thermally couple with the CMB and the resulting fragments are very massive, with $M_{\\rm frag} > 100 M_{\\odot}$ for $z \\ge 5$. Dust-cooling remains relatively insensitive to the presence of the CMB up to metallicities of $Z \\sim 10^{-3} Z_{\\odot}$ and sub-solar mass fragments are formed at any redshift. Above this threshold, heating of dust grains by the CMB at $z \\ge 5$ favors the formation of larger masses, which become super-solar when $Z \\ge 10^{-2} Z_{\\odot}$ and $z \\ge 10$.  \\item Clouds enriched by metals and dust grains with a total metallicity $Z \\ge 10^{-1} Z_{\\odot}$ collapsing at $z > 5$, experience only a single fragmentation epoch (line-induced and dust-induced phases merge) and the characteristic masses are controlled by the CMB temperature, with $M_{\\rm frag} = 0.5, 3, 20 M_{\\odot}$ for $z = 10, 20, 30$.  \\end{enumerate}  The characteristic fragmentation masses are schematically summarized in Fig.~\\ref{fig:summary}, where we report the typical scales associated to dust-induced fragmentation (left panel) and in the absence of dust (right panel) in the different redshift-metallicity regimes discussed above. We have not considered the fragmentation masses associated to line-cooling in the models with dust as they are always $> 50 M_{\\odot}$ independent of the initial metallicity and redshift (see the upper branch in the left panel of Fig.~\\ref{fig:mfrag}).   \\begin{figure*} \\center{{ \\epsfig{figure=summary_dust.eps,height=8.0cm} \\hspace{1.0cm} \\epsfig{figure=summary_nodust.eps,height=8.0cm} }} \\caption{Schematic represtentation of the characteristic masses associated to the relevant fragmentation regimes in the metallicity-redshift plane. In the left panel, we report only the typical scales associated to dust-induced fragmentation, since line-induced fragmentation leads to masses $> 50 M_{\\odot}$ independent of the initial metallicity or redshift. In the right panel, we show the scales expected in models with no dust (see text).} \\label{fig:summary} \\end{figure*}  Our results confirm the general trend suggested by previous semi-analytic (Omukai et al. 2005) and numerical models (Smith et al. 2009), that the CMB tends to inhibit fragmentation of relatively metal-enriched clouds collapsing at high redshift. However, the interplay between metals, dust and the CMB generates a complex dependence of the fragmentation mass on environmental conditions.  It is important to stress that in order to explore the observational consequences of such an environmental dependence, it would be necessary to infer the variation of the functional form of the IMF and not only of the characteristic mass. This issue has been recently addressed by Clark et al. (2009) who predict that, in star forming regions with $Z > 10^{-5} Z_{\\odot}$, competitive accretion can provide a natural explanation for the power-law form of the IMF, which is observed to be remarkably uniform in a wide variety of environments. The competitive accretion of gas from a common reservoir by a collection of protostellar cores requires that the characteristic fragmentation mass be much smaller than the mass of the cooling gas available. As a consequence, Clark et al. (2009) suggest that there may be a link between the CMB and the slope of the IMF and that, in the absence of dust, competitive accretion will not occurr in regions with metallicities $Z \\ge 10^{-3.5} Z_{\\odot}$ collapsing at high redshift whereas it will occurr at all redshifts and metallicities if dust cooling can operate.  As we have mentioned in section 1, one of the main motivation for the present study was to assess whether the CMB could be a viable way to modulate the stellar IMF with redshift as suggested by recent analyses of the elemental abundances of VMP and CEMP stars (Tumlinson 2007; Komiya et al. 2007; Heger \\& Woosley 2008). In Fig.~\\ref{fig:mfrag_red} we show the redshift dependence of the fragmentation mass scales found in the present analysis selecting only the parameter space which is compatible with fragment masses in the range $M_{\\rm frag} \\le 10 M_{\\odot}$ in the presence of dust grains. As a comparison, the thin dotted line shows the empirical dependence of the characteristic stellar mass on redshift proposed by Tumlinson (2007, see eq.~1). It is important to stress that the relation between fragmentation mass scale and characteristic mass shaping the stellar IMF is not straighforward and several processes, such as competitive accretion or dynamical encounters, can act to make it not monothonic. However, the inspection of Fig.~\\ref{fig:mfrag_red} suggests that for all but the highest metallicity, the redshift dependence predicted by our analysis is generally shallower than what inferred from the empirical IMF. Moreover, observations indicate that the fraction of CEMP stars decreases with metallicity being 100\\% at [Fe/H] $<-4.5$, $\\sim 40 \\%$ at [Fe/H] $< -3.5$ and $\\sim 20 \\%$ at [Fe/H] $< -2$ (Beers \\& Christlieb 2005). If CEMP stars reflect the underlying IMF in the range $1 M_{\\odot} < M < 8 M_{\\odot}$,  our results seem to indicate an opposite trend: the effects of the CMB on the characteristic mass increase with metallicity, suggesting that the CMB can not be the only physical driver for the required modulation of the stellar characteristic mass.   \\begin{figure} \\center{{\\epsfig{figure=mfrag_red.eps, height=8.5cm}}} \\caption{Redshift evolution of the fragmentation mass  $M_{\\rm frag}$ for clouds with initial metallicities [M/H] = $-4, -3, -2, -1$ (from bottom to top). The thin dotted line shows the empirical dependence of the stellar characteristic mass on the CMB temperature proposed by Tumlinson (2007, see text).} \\label{fig:mfrag_red} \\end{figure}"
    },
    "0910/0910.3209_arXiv.txt": {
        "abstract": "In the light of the question whether most early-type dwarf (dE) galaxies in clusters formed through infall and transformation of late-type progenitors, we search for an imprint of such an infall history in the oldest, most centrally concentrated dE subclass of the Virgo cluster: the nucleated dEs that show no signatures of disks or central residual star formation. We select dEs in a (projected) region around the central elliptical galaxies, and subdivide them by their line-of-sight velocity into fast-moving and slow-moving ones. These subsamples turn out to have significantly different shapes: while the fast dEs are relatively flat objects, the slow dEs are nearly round. Likewise, when subdividing the central dEs by their projected axial ratio into flat and round ones, their distributions of line-of-sight velocities differ significantly: the flat dEs have a broad, possibly two-peaked distribution, whereas the round dEs show a narrow single peak. We conclude that the round dEs probably are on circularized orbits, while the flat dEs are still on more eccentric or radial orbits typical for an infalling population. In this picture, the round dEs would have resided in the cluster already for a long time, or would even be a cluster-born species, explaining their nearly circular orbits. They would thus be the first generation of Virgo cluster dEs. Their shape could be caused by dynamical heating through repeated tidal interactions. Further investigations through stellar population measurements and studies of simulated galaxy clusters would be desirable to obtain definite conclusions on their origin. ",
        "introduction": "\\label{sec:intro}  Early-type dwarf galaxies largely outnumber all other galaxy types in rich galaxy clusters \\citep[e.g.,][]{vcc}, while they are less abundant on average in galaxy groups, and almost absent in the field \\citep{gu06}. Among dwarf galaxies, the fraction of those with early-type morphology is significantly larger in dynamically more evolved environments \\citep{trentham02}. Apart from this global trend, early-type dwarfs also show a pronounced relation with local galaxy density within a given cluster or group: \\citet{bin87} found that their number strongly increases with local density, analogous to the correlation for giants \\citep{dre80}, and \\citet{tully08} observed that early-type dwarfs strongly cluster around major E/S0 galaxies. These findings have often been interpreted as indicating that early-type dwarfs are actually formed through environmental processes, thereby transforming late-type into early-type objects. Such mechanisms, mainly gas stripping and tidal effects on structure and star formation, have been investigated in various studies \\citep[e.g.,][]{moo98,vZe04a,sab05,boselli08}, with the main conclusion that one or more of these mechanisms must be responsible for the majority of cluster dEs.  Despite the apparently simple structure of early-type dwarfs, a closer investigation of their characteristics revealed a surprising complexity. Disk signatures, such as bars and weak spiral arms \\citep[cf.][]{jer00a}, were identified in a significant fraction of bright Virgo cluster dEs \\citep[and references therein]{p1}, confirming earlier indications of disks \\citep{san84,bin91}. The distribution of projected axial ratios reveals that these objects are not spheroidal, but are rather shaped like thick disk galaxies \\citep{p3}. Similar shapes, though somewhat thicker, are found for the brighter nonnucleated dEs \\citep{fer89}, as well as for those with blue central colors from ongoing or very recent residual star formation \\citep{vig84,p2}. These dE subclasses do not show strong clustering, but are distributed more like the late-type cluster galaxies \\citep{p3}, lending support to the above formation scenarios.   \\begin{figure*} \\epsscale{1.0} \\plotone{f1s.eps} \\caption{{\\bf Selection process and resulting subsamples.} \\emph{Left:} Positions of all Virgo cluster member galaxies (grey dots), of those 149 dE(N)s for which heliocentric velocities are available (grey circles), and of selected fast-moving and slow-moving dE(N)s (grey triangles and black squares, respectively; see the middle panel). The positions of the major elliptical galaxies are indicated by crosses (M49, M84, M86, M87, counterclockwise from South). The central region (see text) is shown as grey-shaded area. \\emph{Middle, top:} Distribution of heliocentric velocities of all dE(N)s (light grey), central fast-moving dE(N)s (dark grey), and central slow-moving dE(N)s (black). The black vertical lines mark the semi-interquartile range ($\\pm 320$ km/s) of velocities relative to the median ($1228$ km/s), calculated from all dE(N)s. The grey vertical lines mark five semi-interquartile ranges, beyond which objects are excluded. \\emph{Middle, bottom:} Distribution of projected axial ratios for the fast and slow subsamples. \\emph{Right, top:} Distribution of projected axial ratios of central dE(N)s. We subdivide them at an axial ratio of 0.8 into a flat and a round subsample. \\emph{Right, bottom:} Resulting distribution of heliocentric velocities for the two subsamples. } \\label{fig:select} \\end{figure*}  Only the nucleated dEs that do not show disk features or blue cores appear to keep up the image of dEs as spheroidal objects consisting of old stars, concentrated in regions of high local density \\citep{fer89,p3}. They are the dE subclass with the oldest stellar populations \\citep{p4,Paudel2009}, and were concluded to be potential primordial objects in a study of the Coma and Fornax clusters \\citep{rak04}. Indeed, semi-analytic galaxy evolution models tied to N-body simulations of hierarchical structure formation show at least qualitative consistency with the observed scaling relations and colors of early-type dwarf and giant galaxies \\citep{DeRijcke2005,janz08,janz09a}, indicating that a ``cosmological'' formation of dEs from 3-sigma density peaks on top of the cluster potential is plausible as well. In the study presented here, our goal is to assess whether even the centrally concentrated nucleated dEs of the Virgo cluster show some indication of an infall history, linking their formation process to environmental effects. ",
        "conclusions": "\\label{sec:conclusion} We have established that different populations, or generations, of Virgo cluster dE galaxies can be distinguished based on their projected shapes and line-of-sight velocities. This is most likely interpreted with the flatter objects moving on radial or highly eccentric orbits, indicating an infalling galaxy population. In contrast, the round dEs have circularized orbits, characteristic for cluster-born galaxies or very early infall. They might thus be the first generation of Virgo cluster dEs. While more detailed stellar population measurements would be necessary for an unambiguous interpretation, the observational findings themselves constitute an important step towards understanding this most abundant galaxy population in clusters. Our results demonstrate the necessity to perform a similar deconvolution of dE populations in any ongoing and future study of dEs, if we aim at tracing back their evolutionary history along with that of their host clusters."
    },
    "0910/0910.1732_arXiv.txt": {
        "abstract": "{} {Astrophysical parameters (\\textit{age, reddening, distance, core and cluster radii}) of 14 open clusters (OCs) projected close to the Galactic plane are derived with 2MASS photometry. The OCs are Be\\,63, Be\\,84, Cz\\,6, Cz\\,7, Cz\\,12, Ru\\,141, Ru\\,144, Ru\\,172, FSR\\,101, FSR\\,1430, FSR\\,1471, FSR\\,162, FSR\\,178 and FSR\\,198. The OCs Be\\,63, Be\\,84, Ru\\,141, Ru\\,144, and Ru\\,172 are studied in more detail than in previous works, while the others have astrophysical parameters derived for the first time.} {We analyse the colour-magnitude diagrams (CMDs) and stellar radial density profiles (RDPs) built after field-star decontamination and colour-magnitude filtered photometry. Field-star decontamination is applied to uncover the cluster's intrinsic CMD morphology, and colour-magnitude filters are used to isolate stars with a high probability of being cluster members in view of structural analyses.} {The open clusters of the sample are located at  $d_\\odot=1.6-7.1$ kpc from the Sun and at Galactocentric distances $5.5-11.8$ kpc, with age in the range 10 Myr to 1.5 Gyr and reddening $E(B-V)$ in the range $0.19-2.56$ mag. The core and cluster radii are in the range $0.27-1.88$\\,pc and $2.2-11.27$ pc, respectively. Cz\\,6 and FSR\\,198 are the youngest OCs of this sample, with a population of pre-main sequence (PMS) stars, while FSR\\,178 is the oldest cluster.} {} ",
        "introduction": "\\label{sec:int}   Open clusters (OCs) are self-gravitating stellar systems formed along the gas- and dust-rich Galactic plane. They contain from tens to a few thousand stars distributed in an approximately spherical structure of up to a few parsecs in radius. The structure of most OCs can be roughly described by  two subsystems, the dense core, and the sparse halo \\citep[][and references therein]{Bonatto2005}.  Because it is relatively simple to estimate the age and distance of OCs, they have become fundamental probes of Galactic disc properties \\citep{Lynga1982, Janes1994, Friel1995, Bonatto2006a, Piskunov2006, Bica2006}. However, the proximity of most OCs to the plane and  the corresponding high values of reddening and field-star contamination  usually restrict this analysis to the more populous and/or to those located at most a few kpc from the Sun \\citep{Bonatto2006a}.  Detailed analysis of OCs and the derivation of their astrophysical parameters will contribute to future disc studies by unveiling the properties of individual OCs. These parameters, in turn, can help constrain theories of molecular cloud fragmentation, star formation, and dynamical and stellar evolution.  \\begin{figure*} \\centering \\includegraphics[scale=0.55,viewport=0 0 470 460,clip]{fig1a.eps} \\includegraphics[scale=0.55,viewport=0 0 470 460,clip]{fig1b.eps} \\caption[]{Left panel: $10\\arcmin\\times10\\arcmin$ XDSS R image of Cz\\,12. Right panel: $10'\\times10'$ XDSS R image of Be\\,84. Images centred on the optimised coordinates.} \\label{fig:1} \\end{figure*}   \\begin{figure*} \\begin{minipage}[b]{0.50\\linewidth} \\includegraphics[width=\\textwidth]{fig2a.eps} \\end{minipage}\\hfill \\begin{minipage}[b]{0.50\\linewidth} \\includegraphics[width=\\textwidth]{fig2b.eps} \\end{minipage}\\hfill \\caption[]{Left panel: 2MASS K$_S$ image $5\\arcmin\\times5\\arcmin$ of FSR\\,198. Right panel: 2MASS K$_S$ image $15\\arcmin\\times15\\arcmin$ of FSR\\,1430. Images centred on the optimised coordinates. The small circle indicates the cluster central region.} \\label{fig:2} \\end{figure*}   The stellar content of a cluster evolves with time, and internal and external interactions affect the properties of individual clusters. Presently the age distribution of star clusters in the disc of the Galaxy can only be explained if these objects are subjected to disruption timescales of a few times $10^{8}$ yrs \\citep{Oort1957, Wielen1971, Wielen1988, Lamers2004}. Open clusters  experience external perturbations by giant molecular clouds (GMCs)  and by spiral arms and other disc-density perturbations. To understand how OCs evolve, it is important to take the effect of these external perturbations into account \\citep{Gieles2007}.  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{fig3.eps}} \\caption[]{Top panels: stellar surface-density $\\sigma (stars\\,\\rm arcmin^{-2}$) of Cz\\,12, computed for a mesh size of $3\\arcmin\\times3\\arcmin$, centred on the coordinates in Table \\ref{tab1}. Bottom: the corresponding isopleth surfaces. Left: observed (raw) photometry. Right: Decontaminated photometry.} \\label{fig:3} \\end{figure}   \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{fig4.eps}} \\caption[]{Same as Fig. \\ref{fig:3} for  FSR\\,198.} \\label{fig:4} \\end{figure}   Cluster disruption is a gradual process with different mechanisms acting simultaneously. Disruption of OCs due to internal processes are characterised by three distinct phases. These phases and their typical timescales are: (\\textit{i}) infant mortality ($\\sim10^{7}$ yr), (\\textit{ii}) stellar evolution ($\\sim10^{8}$ yr) and (\\textit{iii}) tidal relaxation ($\\sim10^{9}$ yr). During all three phases, there are additional external tidal perturbations from e.g. GMCs and disc-shocking that heat the cluster and speed up the process of disruption. However, these perturbations operate on longer timescales for cluster populations and so are more important for tidal relaxation \\citep{Lamers2004, Lamers2005, Lamers2006}. The combination of these effects results in a time-decreasing cluster mass, until either its complete disruption or a remnant \\citep[][and references therein]{Pavani2007} is left.  \\begin{table*} \\centering {\\footnotesize \\caption{Literature and presently optimised coordinates.} \\label{tab1} \\renewcommand{\\tabcolsep}{3.6mm} \\renewcommand{\\arraystretch}{1.1} \\begin{tabular}{lrrrrrrrrr} \\hline \\hline \\multicolumn{5}{c}{Literature}&\\multicolumn{5}{c}{This paper}\\\\ \\cline{2-5} \\cline{7-10} Cluster&$\\alpha(2000)$&$\\delta(2000)$&$\\ell$&$b$&&$\\alpha(2000)$&$\\delta(2000)$&$\\ell$&$b$\\\\ &(h\\,m\\,s)&$(^{\\circ}\\,^{\\prime}\\,^{\\prime\\prime})$&$(^{\\circ})$&$(^{\\circ})$&&(h\\,m\\,s)&$(^{\\circ}\\,^{\\prime}\\,^{\\prime\\prime})$&$(^{\\circ})$&$(^{\\circ})$ \\\\ \\hline Be\\,63 &02 19 36&63 43 00&132.506&2.49& &02 19 30.8&63 43 43&132.49&2.50\\\\ Be\\,84 &20 04 43&33 54 18&70.924&1.27& &20 04 43&33 54 15&70.92&1.27\\\\ Cz\\,6  &02 02 00&62 50 00&130.887&1.05& &02 01 57&62 50 48&130.87&1.07\\\\ Cz\\,7  &02 02 24&62 15 00&131.159&0.52& &02 03 01&62 15 20&131.16&0.53\\\\ Cz\\,12 &02 39 12&54 55 00&138.079&-4.75& &02 39 25&54 54 55&138.11&-4.74\\\\ Ru\\,141 &18 31 19&-12 19 11&19.69&-1.20& &18 31 23&-12 17 50&19.72&-1.21\\\\ Ru\\,144 &18 33 34&-11 25 00&20.749&-1.27& &18 33 33&-11 25 09&20.74&-1.27\\\\ Ru\\,172 &20 11 34&35 35 59&73.11&1.01& &20 11 39&35 37 30&73.14&1.00\\\\ FSR\\,101 &18 49 14&02 46 06&35.147&1.74& &18 49 14&02 46 06&35.14&1.74\\\\ FSR\\,1430 &08 51 52&-44 15 56&264.65&0.08& &08 51 52&-44 16 14&264.66&0.07\\\\ FSR\\,1471 &09 24 08&-47 20 39&270.72&2.14& &09 24 04&-47 20 56&270.71&2.13\\\\ FSR\\,162 &20 01 32&25 14 06&63.211&-2.75& &20 01 26&25 12 30&63.17&-2.74\\\\ FSR\\,178 &20 13 07&29 07 12&67.877&-2.83& &20 13 8.2&29 07 24&67.88&-2.83\\\\ FSR\\,198 &20 02 24&35 41 19&72.184&2.62& &20 02 27&35 40 31&72.17&2.60\\\\ \\hline \\end{tabular} } \\end{table*}     Probably reflecting the Galactocentric-dependence of most of the disruptive effects, the Galaxy presents a spatial asymmetry in the age distribution of OCs. Indeed, \\citet{van1980} noted that OCs older than $\\gtrsim1$ Gyr tend to be concentrated in the anti-centre, a region with a low density of GMCs. In this sense, the combined effect of tidal field and encounters with GMCs has been invoked to explain the lack of old OCs in the solar neighbourhood \\citep[][and references therein]{Gieles2006}. Near the solar circle most OCs appear to dissolve on a timescale shorter than $\\approx1$ Gyr \\citep{Bergond2001, Bonatto2006a}. In more central parts, interactions with the disc, the enhanced tidal pull of the Galactic bulge, and the high frequency of collisions with GMCs tend to destroy the poorly populated OCs on a timescale of a few $10^{8}$ yr \\citep[e.g.][]{Bergond2001}. \\citet{Maciejewski2007} studied a large sample of open clusters, in general not previously studied, to derive fundamental parameters, similar to the present analysis.   This paper is organised as follows. In Sect. \\ref{sec:targ} we provide general data on the target clusters. In Sect. \\ref{sec:tool} we obtain the 2MASS photometry, introduce the tools, CMDs and field-star decontamination algorithm, and derive fundamental parameters of the OCs candidates. In Sect. \\ref{sec:stru} we discuss the stellar radial density profiles (RDPs), colour-magnitude filters, and derive structural parameters. In Sect. \\ref{rel} we discuss properties of the OCs, and concluding remarks are given in Sect. \\ref{conc}. ",
        "conclusions": "\\label{conc} In the present work, we have derived astrophysical parameters of 14 OCs projected close to the Galactic plane by means of 2MASS CMDs and stellar RDPs. Field-star decontamination is applied to uncover the cluster's intrinsic CMD morphology, and CM filters are used to isolate probable cluster members. That field star decontamination leads to consistent CMDs and RDPs shows that we are dealing with OC, instead of field fluctuations. In particular, the present CMD and RDP analyses indicate that 6 IR objects from \\citet{Froebrich2007}, initially identified as cluster candidates from overdensities, are star clusters (Table \\ref{tab4}).  Our sample contains OCs with ages in the range $10\\pm5$ Myr (Cz\\,6 and FSR\\,198) to $1.5\\pm0.5$ Gyr (FSR\\,178), at distances from the Sun in the range $d_{\\odot}\\thickapprox1.6$ kpc (Ru\\,144) to $d_{\\odot}\\thickapprox7.1$ kpc (FSR\\,162) and Galactocentric distances $R_{GC}\\thickapprox5.5$ kpc for Ru\\,141 to $R_{GC}\\thickapprox11.8$ kpc for Be\\,63.  Be\\,84, Ru\\,141, and Ru\\,144 are relatively young surviving OCs located inside the solar circle. Clusters in that region are expected to suffer important tidal stress in the form of shocks from disc and bulge crossings, as well as encounters with massive molecular clouds. In the long run, these processes tend to dynamically heat a star cluster, which enhances the rate of low-mass star evaporation and produces a cluster expansion on all scales. However, for some clusters, mass segregation and evaporation may also lead to a phase of core contraction. Consequently, these effects tend to disrupt most clusters, especially the less populous ones. On the other hand, FSR\\,1430, FSR\\,1471 and Cz\\,12 are older. One of the reasons for such longevity may be their large Galactocentric distance, which minimises the disruption effects. The newly formed open clusters Cz\\,6 and FSR\\,198 show PMS stars in the CMDs.  The present study contributes new open cluster parameters and some revisions to the DAML02 and WEBDA open cluster databases. \\vspace{0.8cm}  \\textit{Acknowledgements}: We thank an anonymous referee for suggestions. We acknowledge support from the CNPq and CAPES (Brazil)."
    },
    "0910/0910.4784_arXiv.txt": {
        "abstract": "The Kuiper belt object Orcus and its satellite Vanth form an unusual system in the Kuiper belt. While most large Kuiper belt objects have small satellites in circular orbits \\citep{2008ssbn.book..335B} and smaller Kuiper belt objects and their satellites tend to be much closer in size \\citep{2008ssbn.book..345N}, Orcus sits in between. Orcus is amongst the largest objects known in the Kuiper belt, but the relative size of Vanth is much larger than that of the tiny satellites of the other large objects \\citep{2008ssbn.book..335B}. Here we characterize the physical and orbital characteristics of the Orcus-Vanth system in an attempt to distinguish discuss possible formation scenarios. From {\\it Hubble Space Telescope} observations we find that Orcus and Vanth have different visible colors and that Vanth does not share the water ice absorption feature seen in the infrared spectrum of Orcus. We also find that Vanth has a nearly face-on circular orbit with a period of $9.5393\\pm0.0001$ days and semimajor axis of $8980\\pm20 $km, implying a system mass of $6.32\\pm 0.01 \\times 10^{20}$ kg or 3.8\\% the mass of dwarf planet Eris. From {\\it Spitzer Space Telescope} observations we find that the thermal emission is consistent with a single body with diameter $940\\pm70$ km and a geometric albedo of $0.28\\pm0.04$. Assuming equal densities and albedos, this measurements implies sizes of Orcus and Vanth of 900 and 280 km, respectively, and a mass ratio of 33. Assuming a factor of 2 lower albedo for the non-icy Vanth, however, implies sizes of 820 and 640 km and a mass ratio of 2. The measured density depends on the assumed albedo ratio of the two objects but is approximately $1.5\\pm0.3$ g cm$^{-3}$, midway between typical densities measured for larger and for smaller objects. The orbit and mass ratio is consistent with formation from a giant impact and subsequent outward tidal evolution and even consistent with the system having now achieved a double synchronous state. Because of the large angle between the plane of the heliocentric orbit of Orcus and the plane of the orbit of Vanth, the system can equally well be explained, however, by initial eccentric capture, Kozai cycling which increases the eccentricity and decreases the pericenter of the orbit of Vanth, and subsequent tidal evolution inward. We discuss implications of these formation mechanisms. ",
        "introduction": "Orcus is among the brightest known bodies in Kuiper belt and appears unique in several ways. Its reflectance spectrum shows the deepest water ice absorption of any Kuiper belt object (KBO) that is not associated with the Haumea collisional family \\citep{2005A&A...437.1115D,2005ApJ...627.1057T,2008A&A...479L..13B}. Compositionally, Orcus appears in many ways to be at a transitional size between the moderately common medium-sized Kuiper belt objects (KBOs) which tend to be spectrally bland \\citep{2008AJ....135...55B,2009Icar..201..272G} and the rare large Kuiper belt objects which are massive enough to retain volatiles such as methane on their surfaces \\citep{2007ApJ...659L..61S}. In addition, Orcus has one of the brightest satellites among the large KBOs \\citep{2008ssbn.book..335B}, relative to the primary. The large fraction of small satellites around large KBOs led Brown et al. (2005) to suggest that many large KBOs suffered collisions which led to satellite formation. In contrast, the similar brightnesses and large angular momenta of the binary members of other KBO systems has been used to suggest a capture origin for these objects \\citep{2002Natur.420..643G,2008ssbn.book..345N}. Again, Orcus appears to be near the boundary of these two regimes.  Another dichotomy exists between the larger and the smaller KBOs that have satellites. While Eris, Pluto, and Haumea -- which all have small presumably collisionally formed satellites --  have densities $\\sim$2 g cm$^{-3}$ or higher \\citep{2007Sci...316.1585B, 2006AJ....132..290B,2006ApJ...639.1238R}, smaller Kuiper belt binaries have densities of 1 g cm$^{-3}$ and even lower \\citep{2006ApJ...643..556S, 2008Icar..197..260G}. Orcus, being intermediate in size between the two populations, could help to shed insight on this apparent bifurcation.  In order to investigate the properties of the Orcus system, we obtained Hubble Space Telescope (HST) observations to determine the orbit, and optical and infrared colors of the components of the system, and Spitzer Space Telescope observations of thermal emission in order to determine the size and thus density of the system. ",
        "conclusions": "Two dynamical scenarios appear equally plausible for explaining the formation of the Orcus-Vanth binary. The orbital characteristics of Vanth are consitent with formation in a sub-catastrophic giant impact such as that that is thought to have formed the Pluto-Charon system\\citet{2005Sci...307..546C}. Indeed, with Orcus in a nearly identical heliocentric orbit as Pluto suggesting an analogy between Pluto-Charon formation and Orcus-Vanth formation is pleasing.  The Orcus-Vanth system could be equally well explained, however, by capture, Kozai cycling, and subsequent tidal evolution (which {\\it cannot} explain Pluto-Charon due to the presence of additional circular coplanar satellites in that system). While the 2.54 magnitude brightness difference between Orcus and Vanth is moderately extreme compared to other binaries which are thought to have formed by capture \\citep{2008ssbn.book..345N}, if Vanth has an albedo lower than Orcus by a factor of two, as seems plausible, the equivalent magnitude difference is only 1.8, closer to values observed in typical systems.  We have hypothesized that perhaps all collisionally formed satellites should share similar spectral characteristics, with the neutral colors and deep water ice absorptions found in the Haumea system and on Charon. While the presence of these unusual spectral features on Vanth would have been significant evidence for a collisional origin it is difficult to argue that their absence rules out collisional origin. Without a significantly more detailed understanding of the physics of icy body collsions and subsequent surface evolution we are unwilling to assert that collision could not have formed a satellite with a surface like Vanth.  Finally, we have noted that the largest KBOs with small presumably collisionally formed satellites all have densities higher than those of the smaller satellites. The uncertainty in the density measured for Orcus, unfortunately, perfectly straddles the range between those low density captured objects and higher density objects like Pluto and Charon (though it is inconsistent with the extremely high densities of objects like Haumea and Quaoar).  The origin of the Orcus-Vanth binary remains uncertain. Future observations which may help to constain this origin include more precise measurements of the size of Orcus, to determine whether Orcus fits into the high or into the low densities found in the Kuiper belt (or elsewhere), and a deep search for additional coplanar satellites, analogous to Nix and Hydra, which would firmly rule out a capture origin.   {\\it Acknowledgments:} This research has been supported by grants from STScI and SSC and through the NASA Earth and Space Science Fellowship program."
    },
    "0910/0910.4098_arXiv.txt": {
        "abstract": "High redshift Type Ia supernovae (SNe~Ia) are likely to be gravitationally lensed by dark matter haloes of galaxies in the foreground. Since SNe~Ia have very small dispersion after light curve shape and colour corrections, their brightness can be used to measure properties of the dark matter haloes via gravitational magnification. We use observations of galaxies and SNe~Ia within the Great Observatories Origins Deep Survey (GOODS) to measure the relation between galaxy luminosity and dark matter halo mass. The relation we investigate is a scaling law between velocity dispersion and galaxy luminosity in the $B$-band: $\\sigma=\\sigma_*(L/L_*)^{\\eta}$, where $L_*=10^{10}h^{-2}L_{\\sun}$. The best-fitting values to this relation are $\\sigma_*=136$ km s$^{-1}$ and $\\eta=0.27$. We find $\\sigma_* \\la 190$ km s$^{-1}$ at the 95 per cent confidence level. This method provides an independent cross-check of measurements of dark matter halo properties from galaxy--galaxy lensing studies. Our results agree with the galaxy--galaxy lensing results, but have much larger uncertainties. The GOODS sample of SNe Ia is relatively small (we include 24 SNe) and the results therefore depend on individual SNe~Ia. We have investigated a number of potential systematic effects. Light curve fitting, which affects the inferred brightness of the SNe~Ia, appears to be the most important one. Results obtained using different light curve fitting procedures differ at the 68.3 per cent confidence level. ",
        "introduction": "Weak gravitational lensing can lead to stretched or magnified (or de-magnified) images of distant sources. The stretching of an image depends on the shear of the lens, whereas the magnification is related to the convergence of the lens. Measurements of shear through the ellipticity of galaxies lensed by galaxies (galaxy--galaxy lensing) have successfully been used to probe the nature of dark matter haloes \\citep{tys84,bra96,hud98,fis00,guz02,hoe03,hoe04,kle06}. In this paper we use measurements of convergence rather than shear to investigate properties of dark matter haloes.  The convergence can be obtained from the brightness of standard candles, because magnified standard candles are brighter than de-magnified ones.  As standard candles we use SNe~Ia, which after light curve shape and colour corrections have small dispersion. Gravitational lensing adds extra scatter to the Hubble diagram \\citep{kan95,fri97,wam97,hol98} and is consequently regarded as a systematic uncertainty in the context of SN~Ia cosmology. This extra dispersion, which increases with redshift \\citep[see, e.g.,][]{ber00,hol05,gun06}, can also be used as a cosmological probe \\citep{met99,dod06}. The distribution of SN~Ia Hubble diagram residuals, which is assumed to trace the magnification probability distribution, can be used as another gravitational lensing probe \\citep{wan05}, which is sensitive to the fraction of compact objects in the Universe \\citep{rau91,metsil99,sel99,mor01,min02,met07}.  Here we use yet another method which utilises individual SNe~Ia rather than the distribution of residuals. We cross-correlate SN~Ia brightnesses and gravitational magnifications computed using properties of galaxies in the foreground, as suggested by \\citet{met98}. For an alternative, potentially very powerful method, also utilising correlations between foreground galaxies and supernova brightnesses, see \\citet{met01}.  For our investigation we use a sample of SNe~Ia which was obtained within the Great Observatories Origins Deep Survey (GOODS). Owing to the {\\it HST} observations, which were part of the GOODS, the sample contain some of the most distant SNe~Ia ever observed \\citep{rie04,rie07}. The high redshift of the supernovae means that they are likely to be significantly magnified by the matter in the foreground \\citep{ber00,hol05}. This magnification is then used to infer the properties of the lensing galaxies.  The outline of the paper is as follows. In section~\\ref{sec:data} we describe the sample of SNe~Ia and the galaxy catalogues we use. The method is explained in section~\\ref{sec:method}. The results obtained when applying the method to the data are presented in section~\\ref{sec:sisresult}. Systematic effects which could affect the results are investigated in section~\\ref{sec:sys}. Section~\\ref{sec:union} is devoted to a particularly important systematic effect, namely the differences in the results which arise due to different light curve fitting packages. The results are finally summarised and discussed in section~\\ref{sec:disc}.  Throughout the paper, we assume a flat universe dominated by dark matter and a cosmological constant. Unless otherwise stated $\\Omega_{\\rm M}=0.3$ and $\\Omega_{\\Lambda}=1-\\Omega_{\\rm M}$. ",
        "conclusions": "\\label{sec:disc} Distant SNe~Ia at high redshift might be affected by gravitational lensing due to matter in the foreground. Gravitational lensing magnification or de-magnification should correlate with Hubble diagram residuals. By exploiting this correlation \\citep{jon07}, we measure the relation between galaxy luminosity and the velocity dispersion of dark matter haloes using data from the GOODS.  For a scaling law, $\\sigma=\\sigma_*(L/L_*)^{\\eta}$, describing the relationship between galaxy absolute $B$-band luminosity, $L$, and dark matter velocity dispersion, we find the best-fitting parameters to be $\\sigma_*=136$ km s$^{-1}$ and $\\eta=0.27$ for $L_*=10^{10}h^{-2}L_{\\sun}$. These results agree well with the results from galaxy--galaxy lensing studies. In a galaxy--galaxy lensing study \\citet{kle06} found $\\sigma_*=132^{+18}_{-24}$ km s$^{-1}$ and $\\eta=0.37 \\pm 0.15$ for the same scaling law and halo model as we use here. \\citet{hoe04} also used galaxy--galaxy lensing to investigate dark matter, but used the truncated isothermal sphere model \\citep{bra96}, which has an additional parameter describing the truncation radius, $s$, of the halo. For a value of $\\eta$ fixed to 0.3 they found $\\sigma_*=136 \\pm 5 \\pm 3$ km s$^{-1}$ (statistical and systematic uncertainties), which agrees with our result, and $s=185^{+30}_{-28}h^{-1}$ kpc.  The constraints derived from galaxy--galaxy lensing are stronger than the ones derived here. Considering how small our sample of lensed SNe~Ia is, it is encouraging that we can derive any constraints at all. Our results rely upon a different aspect of gravitational lensing and therefore provide an independent cross-check of the galaxy--galaxy lensing results.  We have also investigated different systematic uncertainties. The effects of the uncertainties in the cosmological model, photometric redshifts, and magnitude limit appears to be small. Light curve fitting is another systematic effect, which unlike the other ones cannot be ignored for the data set used here. Different light curve fitting packages yield different distances for the same light curves. The Gold and Union samples, which are both small and include different GOODS SNe~Ia, give rather different results. The Union sample is consistent with the no lensing hypothesis ($\\sigma_*=0)$ at the 68.3 per cent confidence level. The Gold sample is consistent with the no lensing hypothesis at the 95, but not at the 68.3 per cent confidence level. From Fig.~\\ref{fig:chi2bf} it is clear that much of the differences can be attributed to individual SNe~Ia. Clearly, light curve fitting is an issue which has to be resolved before SNe~Ia can be used to their full potential as probes of the dark side of the Universe.   The effect of removing two SNe~Ia from the Gold sample was shown  to be quite large in section~\\ref{sec:dust}. For this reason we have performed a bootstrapping exercise to investigate the robustness of our results. The open circles in Fig.~\\ref{fig:boot} show the best-fitting values for 1000 simulated data sets based on resampling of the Gold sample. To facilitate comparison, the best-fitting values (indicated by the solid circle) and confidence level contours for the original sample are also plotted in the figure. The highest density of points is not found near the best-fitting values of the original sample, but closer to $\\eta \\sim -0.2$ and $30 \\la \\sigma_*\\la 120$ km s$^{-1}$. This is another indication that our results depends on individual SNe~Ia.  From our study it is clear that more high redshift SNe~Ia are needed to obtain meaningful results. We anticipate that the final data from the Supernova Legacy Survey \\citep[SNLS,][]{ast06} will be ideally suited for an investigation of the type presented here. According to \\citet{jon08} a firm detection of the correlation between gravitational magnification and SN~Ia brightness can be expected for the final SNLS data set. The SNLS produces not only a homogeneous set of well calibrated SNe~Ia, but also the necessary deep multi colour observations of the foreground galaxies.  \\begin{figure} \\includegraphics[width=84mm,angle=-90]{f11.eps} \\caption{Result of bootstrapping exercise. Open circles correspond to the best-fitting values of $\\sigma_*$ and $\\eta$ obtained for 1000 simulated data sets obtained by resampling of the Gold sample. For comparison the best-fitting values (indicated by the filled circle) and confidence level contours for the original Gold sample are also shown. } \\label{fig:boot} \\end{figure}"
    },
    "0910/0910.3521.txt": {
        "abstract": "We use H$\\alpha$ and far-ultraviolet (FUV, 1539 \\AA) GALEX data for a large sample of nearby objects to study the high mass ($m$ $\\geq$ 2 M$\\odot$) star formation activity of normal late-type galaxies. The data are corrected for dust attenuation using the most accurate techniques at present available, namely the Balmer decrement for H$\\alpha$ data and the total far-infrared to FUV flux ratio for GALEX data. The sample shows a highly dispersed distribution in the H$\\alpha$ to FUV flux ratio (Log $f(H\\alpha)/f(FUV)$ = 1.10 $\\pm$ 0.34 \\AA) indicating that two of the most commonly used star formation tracers give star formation rates with uncertainties up to a factor of 2-3. The high dispersion is partly due to the presence of AGN, where the UV and the H$\\alpha$ emission can be contaminated by nuclear activity, highly inclined galaxies, for which the applied extinction corrections are probably inaccurate, or starburst galaxies, where the stationarity in the star formation history required for transforming H$\\alpha$ and UV luminosities into star formation rates is not satisfied. Excluding these objects, normal late-type galaxies have Log $f(H\\alpha)/f(FUV)$ = 0.94 $\\pm$ 0.16 \\AA, which corresponds to an uncertainty of $\\sim$ 50\\% on the SFR. The H$\\alpha$ to FUV flux ratio of the observed galaxies increases with their total stellar mass. If limited to normal star forming galaxies, however, this relationship reduces to a weak trend that might be totally removed using different extinction correction recipes. In these objects the H$\\alpha$ to FUV flux ratio seems also barely related with the FUV-H colour, the H band effective surface brightness, the total star formation activity and the gas fraction. The data are consistent with a Kroupa (2001) and Salpeter initial mass function in the high mass stellar range ($m$ $>$ 2 M$\\odot$) and imply, for a Salpeter IMF, that the variations of the slope $\\gamma$ cannot exceed 0.25, from $\\gamma$ = 2.35 for massive galaxies to $\\gamma$ = 2.60 in low luminosity systems. We show however that these observed trends, if real, can be due to the different micro history of star formation in massive galaxies with respect to dwarf systems. ",
        "introduction": "The formation and evolution of galaxies can be observationally constrained through the study of their present and past star formation activity. While the star formation history of galaxies can be well determined by fitting their stellar ultraviolet (UV) to near-infrared spectral energy distribution with population synthesis models, the direct observation of the youngest high mass stars is generally used to infer their ongoing activity. The present day star formation rate (SFR) is generally determined using the luminosity of the youngest stars and the relation $SFR$ = $K(\\lambda)$ $L(\\lambda)$, where $K(\\lambda)$ can be inferred from population synthesis models under several assumptions (see Kennicutt 1998 for a review). Hydrogen Balmer emission lines and UV fluxes, both related with the emission of the youngest stellar populations, are thus often used as direct tracers of the present day star formation activity of galaxies. This widely used technique can be blindly applied only if the initial mass function (IMF) is universal, thus independent of morphological type, luminosity and redshift and if the star formation activity of the target galaxies has been constant for a time $\\geq$ than the lifetime of the emitting stars, i.e. $\\sim$ 10$^7$ years for Balmer emission lines and a few 10$^8$ years for the UV ($\\lambda$ $<$ 2000 \\AA). \\\\ Direct measurements of the IMF in the Milky Way or in other very nearby galaxies (Massey et al. 1995; Scalo 1998; Selman \\& Melnick 2008) are consistent with a universal IMF (e.g. Scalo 1986; Kroupa 2001; Renzini 2005). A few recent indirect observational results on the mass to light ratio of low surface brightness galaxies (Lee et al. 2004), on the SED properties of SDSS (Hoversten \\& Galzebrook 2008) and starburst (Rieke et al. 1993) galaxies, on the H$\\alpha$ to UV flux ratio of nearby objects (Meurer et al. 2009) as well as on the star formation activity vs. stellar mass assembly in the far universe (Dav\\'e 2008; Wilkins et al. 2008), combined with theoretical (Krumholz \\& McKee 2008) and statistical considerations (Pflamm-Altenburg et al. 2007) led several authors to question this statement. These works suggest that the IMF, in particular when considered as a unique function representative of the whole galaxy, might not be universal but rather changing with the galaxy stellar mass, the star formation activity or the gas column density. The IMF is a key parameter in the process of star formation: proving its non universality would thus have enormous implications on the study of the star formation process at all scales, from the physics of the interstellar medium to the formation and evolution of galaxies since the earliest phases of the universe. Indeed the IMF is a key ingredient in all models of galaxy evolution. \\\\ It should also be tested whether the condition of stationarity in star formation, required for transforming H$\\alpha$ and UV fluxes into star formation rates, is satisfied in all kinds of galaxies. There are indeed some indications that this is not the case: it has been shown that a bursty history of star formation is expected in dwarf galaxies (Fioc \\& Rocca-Volmerange 1999), consistent with the determination of the star formation history obtained using the stellar colour-magnitude diagram of dwarf galaxies in the local group (e.g. Mateo 1998). A recent starburst activity has been also proposed to explain the the H$\\alpha$ to UV flux ratio of UV luminous galaxies (Sullivan et al. 2000).\\\\ %see Theis & Koeppen astroph submitted It is thus time to investigate whether these two assumptions on the constancy of the star formation history of late-type galaxies and on the universality of the IMF are consistent with the latest observational data in our hands. The comparison of the star formation activity determined using independent indicators of a well defined sample of galaxies with a complete dataset of H$\\alpha$, UV and IR imaging and optical ($R$ $\\sim$ 1000) integrated spectroscopy allows us to quantify the uncertainty associated to different tracers. Furthermore since the ratio of the number of ionizing and non-ionizing photons emitted by a galaxy depends on the slope and on the upper mass cutoff of the IMF, on its present and past star formation activity and on its metallicity, the H$\\alpha$ to UV flux ratio (properly corrected for dust attenuation) can be used to constrain the shape of the IMF.\\\\ Determining the UV ionizing and non-ionizing emitted radiation through observations is however extremely difficult since at these wavelengths the emission of galaxies is severely attenuated by the dust of the interstellar medium. Both UV continuum and Balmer line data need thus to be corrected for dust attenuation.  Quantifying the dust attenuation of the UV radiation is still one of the major challenges in modern astronomy. Models and observations have recently shown that dust extinction can be well determined considering that the UV radiation absorbed by dust is re-emitted in the far-IR (Buat \\& Xu 1996; Witt \\& Gordon 2000; Buat et al. 2002; Boselli et al. 2003; Cortese et al. 2008).  At the same time the attenuation of the Balmer lines can be determined using the Balmer decrement (e.g. Lequeux et al. 1981). This is generally done by comparing the observed to the nominal H$\\alpha$ to H$\\beta$ flux ratio, this last being fairly constant with a value of 2.86 (case B, Osterbrock 1989), provided that spectroscopic data have a sufficient resolution for deblending the [NII] from H$\\alpha$ and for accurately measuring the underlying Balmer absorption. The ionizing radiation could also be absorbed by dust or escape the galaxy before ionizing the gas. All these phenomena should be taken into account in the determination of the emitted radiation using observational data.\\\\ The study of the shape of the IMF using H$\\alpha$ and UV data was first proposed by Buat et al. (1987) and later adopted by Meurer et al. (1999), Boselli et al. (2001) and Charlot et al. (2002). Although limited by the statistics of their sample and by the lack of corollary data for accurate extinction corrections, all these works agreed with a constant IMF. The advent of the GALEX mission (Martin et al. 2005), which provided us with homogeneous UV data for thousands of objects, and the availability of high quality integrated spectroscopy for hundreds of galaxies in the nearby universe (Kennicutt 1992; Jansen et al. 2000; Gavazzi et al. 2004; Moustakas \\& Kennicutt 2006) with IR data, are providing us with a unique opportunity to revisit this subject with a new dataset. The work of Salim et al. (2007), based on GALEX data and SDSS spectra, or that of Meurer et al. (2009) suggested that the H$\\alpha$ to UV flux ratio could indeed vary with stellar mass.  \\\\ The aim of the present work is to study the high mass star formation activity of normal, late-type galaxies. The paper is structured as follows: in sect. 2 we give a brief description of the sample, in sect. 3 of the dataset. The analysis and the discussion are given in sect. 4 and 5 respectively. A detailed description of the observational data (A), of the derived parameters (B) used in the analysis with their uncertainties (C) and on possible selection biases (D) are presented in the Appendix. ",
        "conclusions": "We have studied the high mass star formation activity of an optically/near-IR selected sample of late-type galaxies using two independent, direct tracers of star formation, the H$\\alpha$ and the FUV luminosity. This analysis brought the following results:\\\\  1) The distribution of the H$\\alpha$ over FUV flux ratio is highly dispersed (Log $f(H\\alpha)/f(FUV)$ = 1.10 $\\pm$ 0.34 \\AA) when the data are properly corrected for dust attenuation using the Balmer decrement and the TIR to FUV flux ratio. The dispersion increases (Log $f(H\\alpha)/f(FUV)$ = 1.00 $\\pm$ 0.41 \\AA) when are included data corrected using statistical recipes. This result indicates that H$\\alpha$ and UV luminosities, if blindly combined with standard recipes as those given by Kennicutt (1998), give star formation rates with at best an uncertainty of a factor of $\\sim$ 2-3 depending on the applied dust attenuation correction. The error is significantly reduced (Log $f(H\\alpha)/f(FUV)$ = 0.94 $\\pm$ 0.16 \\AA) once AGN, starburst and highly inclined late-type galaxies are excluded.\\\\  2) Highly inclined galaxies, generally characterized by prominent dust lanes, have H$\\alpha$ over FUV flux ratio significantly higher (Log $f(H\\alpha)/f(FUV)$ = 1.44 $\\pm$ 0.32 \\AA) than normal, late-type galaxies. This result suggests that the dust attenuation corrections applied to this particular class of objects is wrong, probably because the assumption of isotropy in the UV emission of the galaxy is not correct (Panuzzo et al. 2003) or because the Balmer decrement does not give an accurate estimate of the extinction (photon leakage and/or porosity in the ISM). \\\\  3) Extinction is the major source of uncertainty in the determination of the H$\\alpha$ to FUV flux ratio even in normal, late-type galaxies. A direct measure of the Balmer decrement and, to a lesser extent, of the TIR/FUV flux ratio are mandatory for avoiding any residual, systematic effect in the determination of the emitted H$\\alpha$ and FUV fluxes, thus on the star formation activity of the target galaxies.\\\\  4) The data are consistent with a Kennicutt (1998), Kroupa (2001) and Salpeter IMF. To match the data, the fraction of ionizing photons absorbed by dust (1-$f$) should be smaller than previously estimated ($f$=0.57, Hirashita et al. 2003), with $f$ ranging from $\\sim$ 1 for a Salpeter IMF to 0.77 for Kennicutt (1998). The star formation dependent IGIMF proposed by Pflamm-Altenburg et al. (2007; 2009) can be hardly constrained because of the low dynamic range in star formation of our sample. We can only state that the proposed IGIMF is rejected (at 1 $\\sigma$ level) in massive galaxies if $f$ $<$ 0.91.\\\\  5) The H$\\alpha$ over FUV flux is barely related with stellar mass, FUV-H colour, H band effective surface brightness, star formation rate and gas fraction in normal, late-type galaxies. The steepness of these relationships increases if a mass/metallicity dependent $f$ factor is assumed. With the present dataset it is impossible to state whether these trends are real or due to a residual in the adopted extinction corrections. If real, these trends can be explained by a small variation in the IMF, with a slope $\\gamma$ $\\sim$ 2.35 in massive systems and $\\gamma$ $\\sim$ 2.60 in dwarfs if $f$=1, while larger variations should be invoked if $f$$<$ 1 in massive galaxies. We have however shown that the different micro history of star formation in giant, high surface brightness galaxies, characterized by induced star formation along their spiral arms, with respect to dwarf, low surface brightness systems where the ionizing and non ionizing UV emission is dominated by sporadically formed, giant HII regions can explain all the observed trends previously described. \\\\    The consequences of these results are major not only in the study of the star formation properties of nearby galaxies but also in a cosmological context as, for example, in the determination of the cosmic variation of the star formation activity of the universe. We indeed give an accurate estimate of the error on the star formation activity of galaxies using standard recipes based on two among the most widely used tracers, Balmer lines and UV luminosities. Variations in the IMF, as those claimed by Meurer et al. (2009) and Hoversten \\& Galzebrook (2008), which would have drastic consequences in the determination of the mass assembly through stellar formation during the evolution of the universe, are no more required to explain the observations. This micro history variation scenario, here proposed to explain the observed trends, should be tested in the study of the HII distribution of different nearby galaxies spanning a large range in morphological type, luminosity and surface brightness. If confirmed, its more direct consequence would be that the star formation activity of dwarf galaxies, although almost constant on long time scales (Sandage 1986; Gavazzi et al 1996; Boselli et al. 2001), is bursty on short times scales. The stationarity condition required for transforming H$\\alpha$ and UV luminosities into star formation rates would thus not be satisfied in these low luminosity systems, for which other techniques such as SED fitting would be required for measuring their activity."
    },
    "0910/0910.2731_arXiv.txt": {
        "abstract": "The characteristic size of early-type galaxies (ETGs) of given stellar mass is observed to increase significantly with cosmic time, from redshift $z\\gsim 2$ to the present. A popular explanation for this size evolution is that ETGs grow through dissipationless (``dry'') mergers, thus becoming less compact. Combining N-body simulations with up-to-date scaling relations of local ETGs, we show that such an explanation is problematic, because dry mergers do not decrease the galaxy stellar-mass surface-density enough to explain the observed size evolution, and also introduce substantial scatter in the scaling relations. Based on our set of simulations, we estimate that major and minor dry mergers increase half-light radius and projected velocity dispersion with stellar mass as $\\Re\\propto \\Mstar^{1.09\\pm0.29}$ and $\\sget\\propto \\Mstar^{0.07\\pm0.11}$, respectively. This implies that: 1) if the high-$z$ ETGs are indeed as dense as estimated, they cannot evolve into present-day ETGs via dry mergers; 2) present-day ETGs cannot have assembled more than $\\sim 45\\%$ of their stellar mass via dry mergers. Alternatively, dry mergers could be reconciled with the observations if there was extreme fine tuning between merger history and galaxy properties, at variance with our assumptions.  Full cosmological simulations will be needed to evaluate whether this fine-tuned solution is acceptable. ",
        "introduction": "In the hierarchical model of structure formation, galaxy mergers are expected to play a central role in the assembly and evolution of galaxies. If present-day early-type galaxies (ETGs) have experienced mergers in relatively recent times ($z\\lsim 1-2$), most of these mergers must have been dissipationless or ``dry'', because the old stellar populations of ETGs \\citep[e.g.][]{Thomas05} are inconsistent with substantial recent star formation.  Present-day ETGs satisfy several empirical correlations between half-light radius $\\Re$, central stellar velocity dispersion $\\sigmazero$, bulge luminosity $L$, stellar mass $\\Mstar$, black-hole mass $\\MBH$ and gravitational lensing mass $\\Mlens$, such as the $L$-$\\sigmazero$ \\citep{Fab76}, $L$-$\\Re$ \\citep{Kor77}, $\\MBH$-$L$ \\citep{Mag98}, $\\MBH$-$\\sigmazero$ \\citep{Ferrarese2000,Geb00}, $\\MBH$-$\\Mstar$ \\citep{Marconi03}, $\\Mlens$-$\\Re$ and $\\Mlens$-$\\sigmazero$ relations~\\citep[][hereafter NTB09]{Nip09}, the Fundamental Plane \\citep[hereafter FP, relating $L$, $\\Re$ and $\\sigmazero$;][]{Dre87,Djo87} and the lensing Mass Plane \\citep[relating $\\Mlens$, $\\sigmazero$ and $\\Re$;][]{Bol07,Bol08b}. The existence of these scaling laws has been used to establish that present-day ETGs did not form by several dry mergers of lower mass progenitors {\\it obeying the same scaling laws}, because as an effect of dry mergers the galaxy size increases rapidly with mass, while the velocity dispersion remains almost constant, in contrast with the observed correlations~\\citep[][NTB09]{Ciotti01,Nip03,Boy06,Ciotti07}.  Photometric and spectroscopic observations of high-redshift ($z\\gsim 1-2$) ETGs suggest that these objects may be remarkably more compact than their local counterparts \\citep[e.g.][hereafter vD08; Saracco, Longhetti, \\& Andreon 2009]{stiav99,Dad05,Tru06,Zir07,Cimatti08,vdW08,vDo08}.  These findings contributed to renew interest in the proposal that dry merging might be a key process in the growth of ETGs, because it is one of the few known mechanisms able to make galaxies less compact \\citep[e.g.][]{Kho06,Hop09a,Naab09,vdW09}.  However, although dry merging works qualitatively in the right direction, it is not clear whether it works quantitatively.  In this Letter we address the question of the viability of dry merging as a process driving the structural evolution of ETGs, by testing whether it is able to produce the local scaling laws with their relatively small intrinsic scatters.  For this purpose, we use the set of numerical simulations of dissipationless minor and major mergers presented in NTB09 and test them against the stellar-mass scaling relations of the Sloan Lenses ACS Survey (SLACS) sample of lens ETGs \\citep{Auger09}. This strategy allows us to construct a direct link between the observed properties of high-redshift ($z\\gsim1-2$) galaxies---for which estimates of the stellar masses are available---and realistic simulated galaxies, in which the relative distribution of luminous and dark matter has been calibrated on the lensing-mass scaling relations (NTB09).   Throughout the paper we adopt \\cite{Sal55} Initial Mass Function (IMF), which is consistent with the gravitational lensing constraints for SLACS ETGs \\citep{Tre09b}. Observational data are converted to this calibration when necessary. The main results of this paper do not depend on this choice, since they involve stellar-mass ratios and the IMF normalization effectively factors out. ",
        "conclusions": "Dry merging makes galaxies less compact. For this reason it has been proposed to be the key to reconcile the compactness of high-$z$ ETGs with the lower stellar-mass surface-density of local ETGs.  Our quantitative analysis based on N-body simulations shows that {\\it it is very hard to reproduce with dry mergers the growth in size of a factor of $\\sim 5$} between $z\\gsim 2$ and today reported in the literature for $\\Mstar\\sim 10^{11}\\Msun$ ETGs, because the average size growth produced by dry merging is not sufficient.  The tightness of the local scaling laws of ETGs provides an additional and stringent constraint for the dry-merging scenario. As we have shown, generic dry mergers tend to increase the scatter rather than decrease it because they cannot preserve simultaneously the correlations between stellar mass, velocity dispersion and effective radius. On the basis of our set of major and minor merging simulations, we conclude that---independent of the actual compactness of higher-$z$ ETGs---{\\it typical present-day massive ETGs did not assemble more than $\\sim 45\\%$ of their stellar mass and grew more than a factor $\\sim 1.9$ in size via dry merging}. A similar upper limit is found also for the total (dark plus luminous) mass assembly on the basis of the lensing-mass scaling laws (NTB09).     Is this then the end of dry mergers as a major evolutionary mechanism? The only way out could be provided by a high degree of fine tuning of the mix of progenitors and orbits. Although not a natural scenario in our opinion, this fine tuning could perhaps reproduce such tight scaling relations and thus save the day.  A systematic exploration of cosmologically motivated merging hierarchies is necessary to find out whether such an extreme fine tuning actually occurs.  The question of what drives the apparent size evolution of ETGs remains open. A possibility is that a combination of different astrophysical processes and observational biases can do the job \\citep[e.g.][]{Hop09c}, but it is not excluded that the answer is to be found exploring more exotic scenarios, such as those assuming non-standard dark matter \\citep{Lee08,Fer09}."
    },
    "0910/0910.2866_arXiv.txt": {
        "abstract": "We present a new on-the-fly (OTF) mapping of CO($J=3-2$) line emission with the Atacama Submillimeter Telescope Experiment (ASTE) toward the $8' \\times 8'$ (or 10.5 $\\times$ 10.5 kpc at the distance of 4.5 Mpc) region of the nearby barred spiral galaxy M~83 at an effective resolution of 25$^{\\prime \\prime}$. Due to its very high sensitivity, our CO($J=3-2$) map can depict not only spiral arm structures but also spur-like substructures extended in inter-arm regions. This spur-like substructures in CO($J=3-2$) emission are well coincident with the distribution of massive star forming regions traced by H$\\alpha$ luminosity and Spitzer/IRAC 8 $\\mu$m emission. We have identified 54 CO($J=3-2$) clumps as Giant Molecular-cloud Associations (GMAs) employing the CLUMPFIND algorithm, and have obtained their sizes, velocity dispersions, virial masses, and CO luminosity masses. We found that the virial parameter $\\alpha$, which is defined as the ratio of the virial mass to the CO luminosity mass, is almost unity for GMAs in spiral arms, whereas there exist some GMAs whose $\\alpha$ are 3 -- 10 in the inter-arm region. We found that GMAs with higher $\\alpha$ tend not to be associated with massive star forming regions, while other virialized GMAs are. Since $\\alpha$ mainly depends on velocity dispersion of the GMA, we suppose the onset of star formation in these unvirialized GMAs with higher $\\alpha$ are suppressed by an increase in internal velocity dispersions of Giant Molecular Clouds within these GMAs due to shear motion. ",
        "introduction": "The dense molecular medium is an indispensable component of our understanding of the star formation law in galaxies. This is because stars are formed from dense molecular cores, not from the diffuse envelopes of Giant Molecular Clouds (GMCs). In order to understand the global distribution of a dense molecular medium in galaxies, we have conducted an extragalactic CO($J=3-2$) imaging survey of nearby spiral galaxies, ADIoS (ASTE Dense gas Imaging of Spiral galaxies). The submillimeter-wave CO($J=3-2$) line emission is a tracer of dense gas, because its Einstein A coefficient is proportional to $\\nu^3$; therefore, the critical density of CO($J=3-2$) line emission is higher than that of CO($J=1-0$) by a factor of $3^3$, i.e., $n_{\\rm H_2}$ $\\sim$ $10^4$ cm$^{-3}$. The current sample galaxies of the ADIoS project includes a GMA in M~31 \\citep{tosaki2007a}, NGC~604 in M~33 \\citep{tosaki2007b}, M~83 \\citep{muraoka2007}, and NGC~986 \\citep[and see a summary therein]{kohno2008}. Several authors also report large-area CO($J=3-2$) mappings covering disk regions of galaxies. For example, \\cite{minamidani2008} performed CO($J=3-2$) line emission observations of 6 GMCs in the Large Magellanic Cloud (LMC). They identified 32 molecular clumps in the GMCs, and obtained their physical properties: kinetic temperature and density of molecular gas. \\cite{walsh2002} carried out the CO($J=3-2$) observation toward $\\sim 6^{\\prime} \\times 4^{\\prime}$ of NGC~6946, and they found that the arm/inter-arm contrast of CO($J=3-2$) is significantly larger than that of CO($J=1-0$). \\cite{wilson2009} performed the CO($J=3-2$) mapping of some members of Virgo Clusters (NGC~4254, NGC~4321, and NGC~4569), and they examined a large-scale gradient in the gas properties traced by CO($J=3-2$)/CO($J=1-0$) intensity ratio.  The Atacama Submillimeter Telescope Experiment \\citep[ASTE:][]{ezawa2004, ezawa2008}, a new project to operate a 10 m telescope in the Atacama Desert of northern Chile, provides us with an ideal opportunity to generate high-sensitivity and large-scale maps of CO($J=3-2$) line emission, because of the high beam efficiency of the telescope, a low system noise temperature due to the excellent receiver system and atmospheric conditions at the site, and its efficient on-the-fly (OTF) mapping capability \\citep{sawada2008}. The excellent performance in CO mapping enables us to detect and resolve a few 100 pc scale structures of molecular gas in nearby galaxy disks. Molecular gas structures on this scale are often referred to as Giant Molecular-cloud Associations (GMAs) \\citep{vogel1988, rand1990}, whose typical masses are $\\sim 10^7 M_{\\odot}$. Observational studies of GMAs in galaxies provide us with invaluable clues to the physics that governs large-scale star formation in the disk regions of galaxies \\citep[e.g.][]{kuno1995, tosaki2003, sakamoto1996, wong2002, lundgren2004}. Of course, it is essential to investigate physical parameters of GMAs themselves, i.e., their distributions, sizes, velocity dispersions, masses, and luminosities in order to understand their origin and nature.  Studies of GMA properties in galaxies have been mainly performed toward the grand-design spiral galaxy M~51. For example, \\cite{rand1990} identified 26 GMAs in the spiral arm and the inter-arm of M~51 at a spatial resolution of 500 pc $\\times$ 350 pc. They found that most GMAs in the spiral arm appeared to be bound, whereas those in the inter-arm appeared to be unbound, although they could only give the upper limit of virial masses for 14 GMAs due to the limitation of sensitivity. \\cite{rand1993} studied the issue of GMA formation and destruction. The author concluded that both gravitational instability and collisional agglomeration of molecular gas were viable mechanisms for producing bound GMAs in the spiral arm. In addition, \\cite{aalto1999} obtained the higher-quality CO($J=1-0$) image of M~51. The authors discussed the velocity gradients across the spiral arms in detail, and they found that the velocity gradients generally become steeper between GMAs than within them. However, these studies are limited to the relatively narrow, central $2^{\\prime}.5$ -- $3^{\\prime}$ region of M~51. So they could not obtain a spatially unbiased GMA sample covering the whole galaxy disk. Moreover, there are few studies of GMA properties toward other galaxies. Therefore, high angular resolution and high sensitivity CO images covering whole galaxy disks are required in order to obtain a sufficient number of GMA samples and to study not only properties of GMAs but also their relationships with star formation.  In this paper, we present new CO($J=3-2$) images of the nearby barred spiral galaxy M~83 using the ASTE, employing OTF mapping mode. M~83 is a nearby, face-on, barred, grand-design spiral galaxy hosting an intense starburst at its center. The distance to M~83 is estimated to be 4.5 Mpc \\citep{thim2003}; therefore, 1$^{\\prime \\prime}$ corresponds to 22 pc. The inclination of M 83 is 24$^{\\circ}$ \\citep{comte1981}. Its proximity and face-on view enable us to resolve not only spiral arm structures but also each GMA within these structures even by single-dish observations at millimeter or submillimeter wavelength. In addition, substructures diverging from the bar-ends and spiral arms are already found in M~83. Figure~1 shows the Spitzer/IRAC 8 $\\mu$m image, depicting not only clear-cut spiral arm structures but also a number of capillary substructures stretching from spiral arms. These substructures are often referred to as ``spurs'' \\citep[and references therein]{wada2008}. In this study, we divided the M~83 disk into two regions; the spiral arm region and the inter-arm region. We refer to the continuous and thick structures seamlessly linked from the bar in the 8 $\\mu$m image as the ``spiral arm'', and refer to the other region, including spurs, as the ``inter-arm''. If we detect molecular gas associated with spurs, we can study details of the difference in GMA properties between spiral arm and inter-arm regions. Until recently, CO($J=3-2$) images of M~83 were limited to the nuclear region and the bar \\citep{petitpas1998, israel2001, dumke2001, sakamoto2004, bayet2006}. \\citet{muraoka2007} presented wider CO($J=3-2$) image including inner spiral arms, but they could not detect CO($J=3-2$) line emission in inter-arm regions and not cover outer spiral arms. In addition, their CO($J=3-2$) map is not sufficiently sensitive to resolve and identify each GMA in the M~83 disk.  The goals of this paper are: (1) to depict the global distribution of CO($J=3-2$) line emission which is wider and more sensitive than previous CO($J=3-2$) image \\citep{muraoka2007}, (2) to identify GMAs throughout the whole M~83 disk, and to obtain their physical parameters (sizes, velocity dispersions, luminosities, and masses), and (3) to investigate differences between GMAs on spiral arms and those in inter-arm regions comparing their physical parameters, and finally to study the origin and the detailed nature of inter-arm GMAs.  Further analysis of the quantitative relationship between CO($J=3-2$)/CO($J=1-0$) ratios and star formation will be presented in forthcoming papers. ",
        "conclusions": "Since our new CO($J=3-2$) data set covers a wider region (10.5 $\\times$ 10.5 kpc) and achieves a significantly low noise level (1 $\\sigma$ $\\sim$ 25 -- 30 mK in $T_{\\rm MB}$), we were able to find many individual CO($J=3-2$) clumps over the M~83 disk. In order to identify each clump objectively and to obtain its physical parameters, we employed the CLUMPFIND algorithm developed by \\citet{williams1994}.  In this section, we discuss relationships among the size, velocity dispersion, and luminosity of each GMA identified as a clump according to the CLUMPFIND algorithm. Then, we study the properties of GMAs in the M~83 disk throughout the comparison of their virial masses and CO luminosity masses. Finally, we investigate the origin and the nature of GMAs in the inter-arm region.  \\subsection{GMA Identification}  First, we ran the CLUMPFIND software setting a peak threshold of 4 $\\sigma$ and a contour increment of 1 $\\sigma$. In this condition, we identified 99 clumps in the whole M~83 disk. Then, we excluded clumps without adjacent pixels exceeding 3 $\\sigma$ or adjacent channel exceeding 3 $\\sigma$. In addition, we excluded clumps whose diameters are smaller than effective spatial resolution of our CO($J=3-2$) data ($25^{\\prime \\prime}$ $\\sim$ 550 pc) after the deconvolution of clump size with 25$^{\\prime \\prime}$ beam, according to the following formula, \\begin{eqnarray} \\theta_{\\rm decon} = \\sqrt[]{\\theta_{\\rm given}^2 - 25^2} \\,\\,\\, . \\end{eqnarray} Here, $\\theta_{\\rm decon}$ is a deconvolved clump size in arcsec unit, and $\\theta_{\\rm given}$ is a clump size determined by the CLUMPFIND algorithm. Finally, we identified 54 clumps in the M~83 disk ($r$ $>$ 1 kpc), and obtained their radii, velocity dispersions, and CO($J=3-2$) luminosities. Hereafter, we refer to these identified clumps simply as ``GMAs'' in this paper. We show an example of the individual spectrum of an identified GMA in figure 8. This shows a Gaussian shape (a fitted line is also displayed in the figure), and the most of the identified clumps tend to have a similar shape. This is consistent with the mist model by \\cite{solomon1987}. We obtained the total luminosity of all the identified GMAs to be 2.1 $\\times$ $10^8$ K km s$^{-1}$ pc$^2$, which corresponds to 44 \\% of total luminosity of the M~83 disk, 4.8 $\\times$ $10^8$ K km s$^{-1}$ pc$^2$. The other 56\\% luminosity may originate in diffuse components of molecular gas.  Note that we were unable to identify narrow line-width GMAs due to finite velocity resolution (5 km s$^{-1}$). The lower limit of line width that we can detect is 10.3 km s$^{-1}$ (corresponding to the velocity dispersion $\\sigma_v$ = 4.4 km s$^{-1}$) for 4 $\\sigma$ peak GMA and 8.2 km s$^{-1}$ ($\\sigma_v$ = 3.5 km s$^{-1}$) for 5 $\\sigma$ peak GMAs if we assume the Gaussian line profile. In this condition, the lower limit of CO intensity and CO luminosity are 1.2 K km s$^{-1}$ and 4.4 $\\times 10^5$ K km s$^{-1}$ pc$^2$, respectively. This corresponds to the lower limit of CO luminosity mass of 1.8 $\\times 10^6$ $M_{\\odot}$ (see subsection 4.3.2).  \\subsection{GMA Distribution in M~83} Figure~9 and 10 show positions of identified GMAs on the CO($J=3-2$) velocity channel maps and on the Spitzer/IRAC 8 $\\mu$m map, respectively. GMAs are widely distributed over the M~83 disk. We divided these GMAs into two groups according to the location: the spiral arm and the inter-arm. We identified 27 GMAs in the spiral arm, and refer to these as ``on-arm GMAs''. Meanwhile, we also identified 27 GMAs in the inter-arm, and refer to these as ``inter-arm GMAs''.  Obtained properties of identified GMAs are summarized in table 3 and 4. Except for some GMAs, $\\sigma_v$ of GMAs are mostly smaller than 10 km s$^{-1}$. Any GMA with larger $\\sigma_v$ exceeding 10 km s$^{-1}$ are distributed near the center and the bar. Strong non-circular motion along the molecular bar \\citep{sakamoto2004, muraoka2009} may affect these GMAs, and consequently increase their velocity dispersions.  \\subsection{Comparison between on-arm GMAs and inter-arm GMAs} \\subsubsection{CO($J=3-2$) luminosity $L^{\\prime}_{CO(3-2)}$}  Comparing properties between on-arm GMAs and inter-arm GMAs, we found the difference in the GMA radius $R$ and the CO($J=3-2$) luminosity $L^{\\prime}_{CO(3-2)}$. Figure~11 shows the histograms of $R$ and $L^{\\prime}_{CO(3-2)}$ for on-arm and inter-arm GMAs. We found that the averaged values of $R$ and $L^{\\prime}_{\\rm CO(3-2)}$ for on-arm GMAs are both greater than those for inter-arm GMAs. In addition, there exist some on-arm GMAs whose $R$ are comparable to 1 kpc and whose $L^{\\prime}_{\\rm CO(3-2)}$ exceeds $10^7$ K km s$^{-1}$ pc$^2$, whereas there is no inter-arm GMA with such great values in $R$ and $L^{\\prime}_{\\rm CO(3-2)}$. This means on-arm GMAs are relatively larger and brighter in CO($J=3-2$) than inter-arm GMAs, suggesting inter-arm GMAs are less massive than on-arm GMAs.   \\subsubsection{Virial Masses and CO Luminosity Masses}  Using physical parameters of GMAs, we estimate their masses in two different ways: CO luminosity masses $M_{\\rm CO}$ and virial masses $M_{\\rm vir}$. The former, $M_{\\rm CO}$, can be calculated from CO($J=1-0$) luminosity $L^{\\prime}_{\\rm CO(1-0)}$ in K km s$^{-1}$ pc$^2$ unit as follows, \\begin{eqnarray} M_{\\rm CO} = 4.1 \\, L^{\\prime}_{\\rm CO(1-0)} \\,\\,\\,\\,  M_{\\odot} \\,\\, . \\end{eqnarray} This is a revised version of the original formula given by \\citet{bolatto2008}, who assumed the $N_{\\rm H_2}$/$I_{\\rm CO}$ conversion factor $X_{\\rm CO} = 2.0 \\times 10^{20}$ cm$^{-2}$$({\\rm K \\,\\, km \\,\\, s^{-1}})^{-1}$. Here, we adopted $X_{\\rm CO} = 1.8 \\times 10^{20}$ cm$^{-2}$$({\\rm K \\,\\, km \\,\\, s^{-1}})^{-1}$ \\citep{dame2001} for M~83 GMAs. However, the rotational transition of our CO data is not $J=1-0$ but $J=3-2$. Then we have to adopt appropriate CO($J=3-2$)/CO($J=1-0$) ratios $R_{3-2/1-0}$. Here, we assumed $R_{3-2/1-0}$ = 0.6 for the M~83 disk according to \\citet{muraoka2007}. Of course, the $R_{3-2/1-0}$ values vary locally; \\cite{muraoka2007} found that $R_{3-2/1-0}$ typically varies in the range of 0.4 to 0.8 in the M~83 disk. Therefore, the errors of $M_{\\rm CO}$ mainly originate in those of adopted $R_{3-2/1-0}$ values and the uncertainty of $X_{\\rm CO}$. The overall error of $M_{\\rm CO}$ was estimated to be about 40\\%.  Then, we estimate $M_{\\rm vir}$ in order to compare it with $M_{\\rm CO}$. $M_{\\rm vir}$ can be calculated as follows \\citep{bolatto2008}. \\begin{eqnarray} M_{\\rm vir} = 1040 \\, \\sigma_v^2 \\, R \\,\\,\\,\\, M_{\\odot} \\end{eqnarray} where $\\sigma_v$ is the velocity dispersion of GMAs in km s$^{-1}$, and $R$ is the radius in parsec. The error of $R$ was estimated to be 80 pc, which corresponds to half of grid spacing of the CO($J=3-2$) map. In addition, the error of line width was estimated to be 2.5 km s$^{-1}$, corresponding to half of the velocity resolution. This yields the error of $\\sigma_v$ to be 1 km s$^{-1}$. The overall error of $M_{\\rm vir}$ was typically $\\pm$30\\%.  Figure~12 shows a relationship between $M_{\\rm CO}$ and $M_{\\rm vir}$ for M~83 GMAs with that for GMAs in M~51 \\citep{rand1990} and GMCs in LMC \\citep{mizuno2001}. We calculate the $M_{\\rm CO}$ in LMC assuming $X_{\\rm CO} = 7 \\times 10^{20}$ cm$^{-2}$$({\\rm K \\,\\, km \\,\\, s^{-1}})^{-1}$ \\citep{fukui2008}. CO masses of on-arm GMAs in M~83 are in the range of $10^7$ to $10^8$ $M_{\\odot}$, which is comparable with those in M~51, $10^7$ to $6 \\times 10^7$ $M_{\\odot}$ \\citep{rand1990}. We found that virial parameter $\\alpha$, which is defined as the ratio of the virial mass to the CO luminosity mass ($M_{\\rm vir}/M_{\\rm CO}$), of on-arm GMAs are almost equal to unity, that is similar to the tendency of GMCs in LMC. This suggests on-arm GMAs are basically in a gravitationally bound state. On the other hand, there exist some inter-arm GMAs with $\\alpha \\sim$ 3 -- 10. This suggests inter-arm GMAs are not in a gravitationally bound state, which is different from on-arm GMAs. Interestingly, a similar situation is also found in M~51. \\cite{rand1990} identified 6 inter-arm GMAs, and $M_{\\rm vir}$ are more than 10 times greater than $M_{\\rm CO}$ for 5 of inter-arm GMAs. On the other hand, $M_{\\rm vir}$ are comparable to $M_{\\rm CO}$ for most on-arm GMAs. We emphasize that our high-quality CO($J=3-2$) data enable us to identify as many as 54 GMAs in the M~83 disk and to confirm the difference in properties between on-arm GMAs and inter-arm GMAs more clearly.  Note that there exists an on-arm GMA with $\\alpha \\sim 10$ (ID num.\\ 27 in table 3) in M~83. This GMA was detected just in 4 $\\sigma$ peak, so this is possibly a ``fake''. In addition, as shown in figure~9, the systemic velocity of this GMA (507.5 km s$^{-1}$) is shifted from the main component (467.5 -- 477.5 km s$^{-1}$) of the spiral arm there. Therefore, we perhaps have to refer to this GMA as an inter-arm GMA if it is not fake but ``true''.   \\subsection{Origin and Nature of Inter-arm GMAs}  Finally, we discuss the origin and the nature of inter-arm GMAs in M~83. We first briefly review the formation process of spurs, which are major substructures in inter-arm regions. According to numerical simulations \\citep{wada2004}, spur structures in inter-arm regions are formed as follows. Molecular gas clumps in spiral arms are first formed by the Kelvin-Helmholtz instability. Then, some clumps fall along the spiral shocks, whereas other gases try to maintain nearly circular rotation. As a result, they move away from the spiral shock towards outer spiral arms. Since a clump has the internal gradient of its angular momentum, the clumps gas eventually extends to inter-arm regions as it rotates in the galactic potential. Therefore, it is suggested that GMAs distributed in inter-arm regions are ensembles of GMCs moving away from spiral arms in this way. These GMCs/GMAs might be extended by shear motion, and finally they could be destroyed and change into atomic phase. However, as can be seen in figure~6, massive star forming regions are distributed even in inter-arm regions, corresponding to CO($J=3-2$) spurs. This suggests dense gas formation and massive star formation can occur within inter-arm GMAs.  Figure~13 shows a correlation between an H$\\alpha$ luminosity and a virial parameter $\\alpha$ for each inter-arm GMA in M~83. An anti-correlation can be seen; GMAs with large $\\alpha$ tend {\\it not} to be associated with massive star forming regions. Then, we investigate what parameter of the GMA determines $\\alpha$. Figure~14 shows a comparison between the radius $R$, velocity dispersion $\\sigma_v$, and $\\alpha$. Obviously, $\\sigma_v$ correlates well with $\\alpha$, whereas little correlation between $R$ and $\\alpha$ can be seen. Therefore, these plots suggest that $\\alpha$ increases depending on $\\sigma_v$ rather than $R$. This dependence is also found for GMCs in the spiral arm of IC~342 \\citep{hirota2008} and at the center of the Milky way \\citep{oka2001}. Considering that a GMA is an ensemble of multiple GMCs and the actual star formation occurs in the dense molecular cores of molecular clouds, it is suggested that increase in the velocity dispersion of a GMC component within a GMA, $\\sigma_{v, {\\rm GMC}}$, prevents the onset of star formation within the GMC and we have observed the increase in $\\sigma_{v, {\\rm GMC}}$ as the increase in the overall $\\sigma_v$ of the GMA due to the inadequate spatial resolution. Of course, the shear motion may also increase GMC-to-GMC velocity dispersions within the GMA, contributing the increase in overall $\\sigma_v$. Therefore, we suggest that there mainly exist two kinds of inter-arm GMAs. One is virialized ($\\alpha$ $\\sim$ 1) GMAs with smaller $\\sigma_v$, which are not affected by the shear motion. Then star formation may occur within these GMAs even after moving away from spiral arms toward the inter-arm region. The other is unvirialized GMAs with larger $\\sigma_v$ and $\\alpha$, which are affected by the shear motion. Star formation process within these GMAs might be suppressed due to the increase in the internal velocity dispersion of GMCs, $\\sigma_{v, {\\rm GMC}}$.  Note that since our CO observations are performed by a single dish telescope, we cannot resolve each GMC within identified GMAs. In order to study the relationship of $\\sigma_{v, {\\rm GMC}}$, $\\sigma_v$, virial mass, and star formation in more detail, high angular resolution observations using radio interferometers are necessary."
    },
    "0910/0910.2261_arXiv.txt": {
        "abstract": "{Studies of the Diffuse Ionized Gas (DIG) have progressed without providing so far any strict criterion to distinguish DIGs from \\hii regions. In this work, we compile the emission line measurements of 29 galaxies that are available in the scientific literature, thereby setting up the first DIG database (\\ddb). Making use of this database, we proceed to analyze the global properties of the DIG using the \\lnii/\\ha, \\loi/\\ha, \\loiii/\\hb and \\lsii/\\ha lines ratios, including the \\ha emission measure. This analysis leads us to conclude that the \\nii/\\ha ratio provides an objective criterion for distinguishing whether an emission region is a DIG or an \\hii region, while the EM(\\ha) is a useful quantity only when the galaxies are considered individually. Finally, we find that the emission regions of Irr galaxies classified as DIG in the literature appear in fact to be much more similar to \\hii regions than to the DIGs of spiral galaxies.\\vspace{0.12cm}}   \\addkeyword{astronomical data bases: miscellaneous} \\addkeyword{galaxies: general} \\addkeyword{ISM: HII regions} \\addkeyword{ISM: lines and bands} \\addkeyword{line: formation}  \\RescaleTitleLengths{1.025}  \\begin{document} ",
        "introduction": "\\label{sec:Intro}   From  observations of free-free continuum absorption, Hoyle \\& Ellis (1963) inferred the existence of a new gas component of the interstellar medium (ISM). They proposed that this new component was warm ($T_e \\sim 10^4$ K), of low density ($n_e \\sim 10^{-1}\\, cm^{-3}$), ionized and that it surrounded the Milky Way's disk. The integrated luminosity of this hypothetical medium was comparable to that expected from photoionization by O and B stars from the disk.   The existence  of this gas component was later confirmed by the signal dispersion from pulsars and by the detection of  faint optical  emission lines from the galactic diffuse interstellar medium (Reynolds, Scherb, \\& Roesler 1973). Deep images of external galaxies such as NGC 891 show a similar warm gas component (Dettmar 1990; Rand 1998). This component has since been found in several other galaxies (e.g., Bland-Hawthorn, Sokolowski, \\& Cecil 1991a,b; Veilleux, Cecil, \\& Bland-Hawthorn 1995; Greenawalt, Walterbos, \\& Braun 1997; Martin \\& Kennicutt 1997; Wang, Heckman, \\& Lehnert 1997). Today, it carries various names, such as WIM (Warm Ionized Medium), DIM (Diffuse Ionized Medium), DIG (Diffuse Ionized Gas) or eDIG (extraplanar DIG), and ``Reynold's Layer'' (in the case of the Galaxy). Hereafter, we will use the generic names DIG and eDIG.  Despite the fact that the DIG in our Galaxy accounts for $\\sim$ 90\\% of the ionized medium's mass (Reynolds 1991), the nature of this low density plasma is far from being well understood. The first problem encountered in DIG studies is how to formally distinguish DIGs from \\hii regions, which  similarly consist of a warm and ionized gas phase. Several methods have been proposed in the literature. Examples are the \\ha equivalent width (Bland-Hawthorn et al. 1991a), the \\ha emission measure\\footnote{The EM in units of pc\\,${\\rm cm^{-6}}$ is defined as $\\mathrm{EM} \\equiv \\int n_e n_{H^+} \\, ds \\approx \\int n_e^2 \\, ds$.}, EM (e.g., Walterbos \\& Braun 1994), or the surface brightness (e.g., Ferguson et al. 1996). Nevertheless, none of the methods has so far been accepted as standard, mainly because the density and size of the emitting regions are usually unknown. To illustrate the above problem, an example is provided by the galaxy M31, for which Walterbos \\& Braun (1994) adopted the EM criterion to distinguish DIGs from \\hii regions. They found that close to low star formation regions, the DIG always had an EM $<50$\\,pc\\,${\\rm cm^{-6}}$, while near higher star formation regions, the latter had an EM near 100\\,pc\\,$\\rm cm^{-6}$. Therefore, even though the EM is proportional to the electron density and should thereby reflect how ``diffuse'' an emission region is, this criterion  was shown to vary with position in this galaxy.  Another problem in the study of the DIG is the difficulty in observing it, mainly because of its low surface brightness. As a result, the observations in the literature are sparse and only a few line ratios are typically reported. Furthermore, systematical comparative studies of the general behavior of the EMs or of line ratios among different objects are still missing. For example, even though a majority of authors agree on the existence of an anti-correlation between the \\lnii/\\ha or \\lsii/\\ha ratios with the EM (these ratios increase when the EM decreases, e.g., Bland-Hawthorn et al. 1991a; Rand 1998), very few studies were done with more than two galaxies (e.g., Wang et al. 1997; Zurita, Rozas, \\& Beckman 2000; T\\\"ullmann \\& Dettmar 2000; Miller \\& Veilleux 2003). In this work, we will verify for the first time whether such an anticorrelation is a general feature of the DIG or not. Another interesting aspect is the apparent confusion concerning the \\loiii/\\hb\\ ratio in the literature, for which some authors report a constant behavior with EM (e.g., Wang et al. 1997, at least for some galaxies), while others report an increase when the EM decreases (e.g., Rand 1998; T\\\"ullmann \\& Dettmar 2000).  The spectral characteristics (different from those of \\hii regions) combined with the large spatial extension of the DIG (typically $|z| \\approx 2$~kpc), are not well reproduced by models that consider only photoionization from O-B stars (e.g., Rand, Wood, \\& Benjamin 2009), even though it could provide the entire UV photon flux required to mantain the DIG ionized (e.g., Reynolds 1990; Zurita et al. 2000). As a result, several authors have suggested that the DIG could be excited by extra ionization (or heating) sources like: shocks (e.g., Sivan, Stasi\\'nska, \\& Lequeux 1986), mixing layers (e.g., Binette et al. 2009), old hot stars (e.g., Sokolowski \\& Bland-Hawthorn 1991), and decaying dark matter (e.g., Dettmar \\& Schultz 1992). In the study of the Seyfert 2 galaxy NGC 1068, Bland-Hawthorn et al. (1991b) proposed photoionization by the AGN power law. None of these models is completely successful in explaining the line ratios, spatial extensions and luminosities of the DIGs in spiral galaxies. For a complete review of the posibles DIG ionizing or heating sources, see Bland-Hawthorn, Freeman, \\& Quinn (1997).  With the previous arguments in mind, and to correct for the lack of comparative studies in the field of diffuse emission, we have compiled a comprehensive database of emission line ratios or fluxes, from DIGs or \\hii regions, published in the literature. This database is hereafter called the Diffuse Ionized Gas Emission Database (\\ddb) which will soon be available at the VizieR website of CDS (Centre de Donn\\'ees Astronomiques de Strasbourg, \\url{http://vizier.u-strasbg.fr}).  The structure of the article is the following: after describing \\ddb in \\S~\\ref{sec:DB}, we proceed with the analysis and interpretation of the data in \\S~\\ref{sec:ID}. A discussion follows in \\S~\\ref{sec:discussion} and a brief conclusion is given in \\S~\\ref{sec:conclusions}.    \\begin{table*}[!t] \\vspace{0.1cm} \\centering \\setlength{\\tabnotewidth}{0.99\\textwidth} \\setlength{\\tabcolsep}{0.77\\tabcolsep} \\tablecols{6} \\caption{Bibliographical references used for the compilation of \\ddb} \\label{referencias} \\scriptsize \\begin{tabular}{cccccc} \\toprule {\\bf Ref\\_ID\\tabnotemark{a}} &{\\bf Author} & {\\bf Galaxy} & {\\bf Slit\\tabnotemark{b}} & {\\bf Observations\\tabnotemark{c} } & {\\bf No. regions} \\\\ \\midrule 1& Benvenuti, D'Odorico, \\& Peimbert (1976) & M33 & Interarm & \\oiii$\\lambda$4959, \\oiii, \\nii & 62\\\\ \\noalign{\\smallskip} 2& Golla, Dettmar, \\& Domgorgen (1996) & NGC 4631 & $\\perp$ & \\nii, \\sii & 39\\\\ & & & $\\parallel$ & \\nii, \\sii & 80\\\\ \\noalign{\\smallskip} 3& Greenawalt et al. (1997) & M31 & DIG, \\hii & \\oii, \\oiii, \\nii, \\siip, EM & 18\\\\ \\noalign{\\smallskip} 4& Martin \\& Kennicutt (1997) & NGC 1569 & DIG & \\oiii, \\he, \\oi, \\nii, \\siip & 25\\\\ \\noalign{\\smallskip} 5& Wang et al. (1997) & M51 & Interarm & \\oiii, \\nii, \\siip & 3\\\\ & & M101 & Interarm & \\oiii, \\nii, \\siip & 6\\\\ & & NGC 2403 & Interarm & \\oiii, \\nii, \\siip & 5\\\\ & & NGC 4395 & Interarm & \\oiii, \\nii, \\siip & 3\\\\ \\noalign{\\smallskip} 6& Rand (1998) & NGC 891 & $\\parallel$ &\\oi, \\nii, \\sii& 45\\\\ & & & $\\perp$ & \\oiii, \\oi, \\nii, \\sii & 19\\\\ \\noalign{\\smallskip} 7& Galarza, Walterbos, \\& Braun (1999) & M31 & DIG, SBRs & \\oii, \\oiii, \\nii, \\siip, EM & 30\\\\ \\noalign{\\smallskip} 8& Hoopes, Walterbos, \\& Rand (1999) & NGC 4631 & $\\perp$ & \\oiii, \\siip, EM& 31\\\\ \\noalign{\\smallskip} 9& T\\\"ullmann \\& Dettmar (2000) & NGC 1963 & $\\perp$ &\\oiii, \\he, \\oi, \\nii, \\sii, EM& 40\\\\ & & NGC 3044 & $\\perp$ &\\oiii, \\he, \\oi, \\nii, \\sii, EM& 47\\\\ & & NGC 4402 & $\\perp$ &\\nii, \\sii, EM& 24\\\\ & & NGC 4634 & $\\perp$ &\\nii, \\sii, EM& 40\\\\ \\noalign{\\smallskip} 10& Collins \\& Rand (2001) & NGC 4302 & $\\perp$ & \\nii, \\siip & 13\\\\ & & NGC 5775 & $\\perp$ & \\oi, \\oiii, \\nii, \\siip & 31\\\\ & & UGC 10288 & $\\perp$ & \\oi, \\oiii, \\nii, \\siip & 19\\\\ \\noalign{\\smallskip} 12& Otte, Gallagher, \\& Reynolds (2002) & NGC 5775 & $\\perp$ & \\oii, \\oiii, \\nii, \\siip & 38\\\\ & & NGC 4634 & $\\perp$ & \\oii, \\oiii, \\nii, \\siip & 34\\\\ & & & $\\parallel$ & \\oii, \\oiii, \\nii, \\siip & 32\\\\ & & NGC 4631 & $\\parallel$ & \\oii, \\oiii, \\nii, \\siip & 34\\\\ & & NGC 3079 & $\\parallel$ & \\oii, \\oiii, \\nii, \\sii & 53\\\\ & & NGC 891 & $\\parallel$ & \\oii, \\nii, \\siip & 50\\\\ \\noalign{\\smallskip} 13& Hoopes \\& Walterbos (2003) & M33 & Interarm & \\oiii, \\he, \\nii, \\sii,  EM&37\\\\ & & M51 & Interarm  & \\oiii, \\nii, \\sii, EM & 26\\\\ & & M81 & Interarm  & \\oiii, \\nii, \\sii     & 2 \\\\ \\noalign{\\smallskip} 14& Miller \\& Veilleux (2003) & UGC 2092 & $\\perp$ &\\nii, \\sii&4\\\\ & & UGC 3326 & $\\perp$ & \\nii, \\sii&6\\\\ & & UGC 4278 & $\\perp$ & \\oiii, \\he, \\oi, \\nii, \\sii&13\\\\ & & NGC 2820 & $\\perp$ & \\oiii, \\he, \\oi, \\nii, \\sii&13\\\\ & & NGC 3628 & $\\perp$ & \\oiii, \\nii, \\sii&13\\\\ & & NGC 4013 & $\\perp$ & \\oiii, \\he, \\oi, \\nii, \\sii&9\\\\ & & NGC 4217 & $\\perp$ & \\oiii, \\oi, \\nii, \\sii&17\\\\ & & NGC 4302 & $\\perp$ & \\oiii, \\he, \\oi, \\nii, \\sii&9\\\\ & & NGC 5777 & $\\perp$ & \\oiii, \\nii, \\sii&6\\\\ \\noalign{\\smallskip} 15& Hidalgo-G\\'amez (2006) & ESO245-G05 & DIG, \\hii & \\oii, \\oiii, \\nii, \\sii&7\\\\ & & Gr8 & DIG, \\hii & \\oiii, \\he, \\nii, \\sii&19\\\\ \\noalign{\\smallskip} 16& Voges \\& Walterbos (2006) & M33 & DIG &\\oiii, \\oi, \\nii, \\sii & 7\\\\ \\noalign{\\smallskip} 17& Voges (2006) & NGC 891 & $\\parallel$ &\\oiii, \\oi, \\nii, \\sii, EM& 16\\\\ & & & $\\perp$ &\\oiii, \\oi, \\nii, \\sii, EM& 10\\\\ \\noalign{\\smallskip} 18& Hidalgo-G\\'amez (2007) & DDO53 & DIG, \\hii &\\oiii, \\sii& 26\\\\ \\bottomrule \\noalign{\\smallskip} \\tabnotetext{a}{These numbers also serve as bibliographical entries in \\ddb.} \\tabnotetext{b}{Slit orientation in each galaxy. The symbols and descriptors are described in \\S~\\ref{sec:descr}.} \\noalign{\\smallskip} \\tabnotetext{c}{The line wavelengths are the following: \\loii, \\loiii, \\lhe, \\loi, \\lnii, \\lsii  and \\lsiimas.} \\end{tabular} \\vspace{0.2cm} \\end{table*} ",
        "conclusions": "\\label{sec:conclusions}  In this work, we present the first emission line database from DIGs (\\ddb) made up from spectroscopic data available in the literature. The database is a compilation of 17 bibliographical references. It contains 1061 observed regions (309 \\hii regions, 218 transition zones and 509 DIGs) out of 29 galaxies (spirals and irregulars) and is freely available from CDS.  \\ddb has allowed us to carry out for the first time a statistical analysis that aimed at characterizing the general behavior of the strongest emission lines, and at finding global trends among the various line ratios, including the EM. One of the results is that the DIGs of irregular galaxies show extreme values in their line ratios with respect to that of spirals (see \\S~\\ref{sec:ID} and \\S~\\ref{sec:discussion}). These, in fact, lie close to the values observed in  \\hii regions. Since the available data are very limited, it is indispensable to obtain more observations to confirm this result.  The analysis carried out with \\ddb leads us to define a universal criterion for spiral galaxies, which allows distinguishing DIGs from \\hii regions. The distribution of \\nii/\\ha ratios shows that \\hbox{\\nii/\\ha $= -0.4$} defines a critical value: any emission region with a ratio greater than $-0.3$ should be classified as DIG while a ratio below $-0.5$ is likely to come from an \\hii region.  Taking advantage of the high number of data in \\ddb, we could confirm (or refute) the reality of line ratio behaviors that have been proposed in the literature concerning individual galaxies. We could check to what extent these applied to all DIGs taken together.  This is the case for the \\nii/\\ha ratio, which shows a well defined anti-correlation with EM, but of varying slopes between galaxies. These variations in the log-log slopes rule out any definition of a universal EM value to distinguish DIGs from \\hii regions. Nevertheless, a critical value, EM$_{\\rm c}$, can be defined for each galaxy individualy.  \\vspace{0.15cm}"
    },
    "0910/0910.2920_arXiv.txt": {
        "abstract": "We present a timing analysis of the 2009 outburst of the accreting millisecond X-ray pulsar Swift J1756.9-2508, and a re-analysis of the 2007 outburst. The source shows a short recurrence time of only $\\sim 2$ years between outbursts. Thanks to the approximately 2 year long baseline of data, we can constrain the magnetic field of the neutron star to be $0.4\\times 10^{8}$G$\\simless B\\simless 9\\times 10^{8}$G, which is within the range of typical accreting millisecond pulsars. The 2009 timing analysis allows us to put constraints on the accretion torque: the spin frequency derivative within the outburst has an upper limit of $|\\dot{\\nu}|\\simless 3\\times 10^{-13}\\rm\\,Hz\\,s^{-1}$ at the 95\\% confidence level. A study of pulse profiles and their evolution during the outburst is analyzed, suggesting a systematic change of shape that depends on the outburst phase. ",
        "introduction": "The accreting millisecond X-ray pulsar (AMXP) Swift J1756.9-2508 (henceforth referred to as J1756) was discovered in 2007 by the Burst Alert Telescope (BAT) aboard the \\textit{Swift} satellite \\citep{kri07a}. Follow up observations were performed with the \\textit{RXTE} Proportional Counter Array (PCA) which revealed an X-ray pulsar with a spin frequency of 182 Hz orbiting in 54.7 min around a very low mass companion of minimum mass 0.0067$\\msun$ \\citep{kri07b}.  The X-ray lightcurve of the 2007 outburst showed a very fast rise of less than 1 day, followed by an exponential decay with an e-folding time of $7.4\\pm 0.4$ days. An inspection of archival \\textit{RXTE} All-Sky Monitor (ASM) data revealed a lack of detections in the previous 10 years of observations, which suggested a long recurrence time for the outbursts. However, due to the low flux of the source and the poor sensitivity of the ASM, the recurrence time was only tentative and a shorter duty cycle was not firmly excluded.  In July 2009 a new outburst was discovered by both \\textit{Swift} BAT and \\textit{RXTE} PCA observations \\citep{pat09}. Follow-up observations with the \\textit{Swift} X-ray telescope (XRT) and  \\textit{RXTE} PCA were promptly scheduled and monitored the outburst of J1756 until it returned into quiescence.  In this Letter we report a first analysis of the 2009 outburst and a re-analysis of the 2007 data to evaluate the long-term evolution of J1756 on a baseline of approximately 2 years. ",
        "conclusions": "Swift J1756.9-2508 has shown a short recurrence time of only two years and it is the optimal candidate to study the long-term evolution of neutron star spin and orbital parameters.  The pulse profiles evolve in time with strong similarities between the two outbursts and show a steep increase of pulsed fraction with energy. With the increasing number of AMXPs exhibiting an hard energy dependence of pulse amplitudes, it becomes more likely to discover new AMXPs among the large family of the non pulsating low mass X-ray binaries. Indeed, all the 12 AMXPs discovered until now have been observed with X-ray telescopes which are sensitive to energies below 20 keV, or that are not designed to perform efficient coherent X-ray timing observations. With the forthcoming X-ray telescopes like ASTROSAT, which will have a very high effective area up to $\\sim80$ keV, the detection of AMXPs with rising pulsed fractions at hard energies could be much easier and efficient.  The pulse phases show a low, but significant, timing noise content and special techniques have to be applied to determine the neutron star spin evolution. Thanks to the measurement of timing parameters in two outbursts, we can now set the first constraints on the neutron star magnetic field ($0.4\\times 10^{8}G\\simless B\\simless 9\\times 10^{8}$G) that makes this neutron star a typical member of the AMXP family."
    },
    "0910/0910.3943_arXiv.txt": {
        "abstract": "{} {We study the role of the quasi-linear diffusion (QLD) in producing $X$-ray emission by means of ultra-relativistic electrons in AGN magnetospheric flows.} {We examined two regions: (a) an area close to the black hole and (b) the outer magnetosphere. The synchrotron emission has been studied for ultra-relativistic electrons and was shown that the QLD generates the soft and hard $X$-rays, close to the black hole and on the light cylinder scales respectively. } {By considering the cyclotron instability, we show that despite the short synchrotron cooling timescales, the cyclotron modes excite transverse and longitudinal-transversal waves. On the other hand, it is demonstrated that the synchrotron reaction force and a force responsible for the conservation of the adiabatic invariant tend to decrease the pitch angles, whereas the diffusion, that pushes back on electrons by means of the aforementioned waves, tends to increase the pitch angles. By examining the quasi-stationary state, we investigate a regime in which these two processes are balanced and a non-vanishing value of pitch angles is created.} {} ",
        "introduction": "One of the major problems related to active galactic nuclei is the origin of the nonthermal high energy radiation. According to standard approaches, the most commonly encountered radiation mechanisms at a level sufficient for application to AGN  is the synchrotron mechanism and the inverse Compton scattering (\\cite{blan}). Because of strong synchrotron losses, relativistic electrons in general, quickly lose their perpendicular energy, and on a synchrotron cooling timescale $\\sim 10^{-6}s-10^{-3}s$, the particles transit to their ground Landau state. In this case, the electrons may be described approximately as moving one-dimensionally along the field lines and the synchrotron radiation to have been absorbed. This is why the broadband emission spectrum of AGN consists of two components: the high-energy (from X-rays to $\\gamma$-rays) component is formed by the inverse Compton scattering and not by the synchrotron mechanism, which is supposed to be involved only in the low-energy (from radio to optical/UV) band. However, under certain conditions, due to the QLD of cyclotron waves, the pitch angles might increase, leading to the efficient production of synchrotron radiation. %  The QLD was applied to pulsars in a series of papers (\\cite{machus1,lomin,malmach}). Malov \\& Machabeli (2001) studied optical synchrotron emission of radio pulsars. In the outer parts of pulsar magnetospheres, these authors demonstrated that because of the cyclotron instability, the transverse momenta of relativistic particles is non-zero, giving rise to the pitch angle distribution, which in turn, via the QLD, leads to the synchrotron emission. Applying the kinetic approach to a particular pulsar, RX J1856.5-3754, \\cite{nino} showed that waves excited by the cyclotron mechanism, in terms of the creation of the pitch angles, come into the radio domain. The QLD interesting because the recent detection of very high energy (VHE) pulsed emission form the Crab pulsar (\\cite{magic}). The MAGIC Cherenkov telescope discovered the pulsed emission above 25GeV between 2007 October and 2008 February. It has been shown that the corresponding VHE signal peaks at the same phase as the signal in the optical spectrum (\\cite{magic}). In turn this indicates that the polar cap models must be excluded from the possible scenario of the radiation. On the other hand, analysis of the MAGIC data implies that the location of the aforementioned VHE and optical radiation must be the same. According to the quasi-linear diffusion, on length scales typical of the light cylinder (a hypothetical zone, where the linear velocity of rigid rotation equals exactly the speed of light), the cyclotron instability occurs in the optical band, leading to an increase in the pitch angles via the QLD. This mechanism automatically explains the coincidence of phases in the optical and VHE bands (\\cite{difus}).  AGN magnetospheres are supported by strong magnetic fields and therefore, the QLD might also be of great importance to these particular objects. As aforementioned, for ultra-relativistic electrons the synchrotron losses are so efficient that the synchrotron mechanism takes place only for relatively low energy particles and highly relativistic electrons are involved in radiation via the inverse Compton scattering. This is not true for the QLD, because as for the pulsar magnetospheres, AGN magnetospheric particles will undergo the QLD, preventing the rapid damping of pitch angles, giving rise to the emission process.  In the present paper, we study the role of the QLD in producing the X-rays via the synchrotron mechanism in AGN magnetospheres. The paper is organized as follows. In Section 2, we consider the kinetic approach to the quasi-linear diffusion, in Sect. 3 we present our results and in Sect. 4 we summarize them. ",
        "conclusions": "\\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics[angle=0]{13515fg1.eps}} \\caption{The synchrotron emission energy versus the cyclotron frequency. The set of parameters is $M_{BH} = 10^8M_{\\odot}$, $n_b = 2000cm^{-3}$, $B = 10^4G$, $\\gamma_p = 2$, $\\gamma_b\\in\\{4\\times 10^7;10^9\\}$, and $\\rho = R_g$. As we see from the plot, the low frequency cyclotron mechanism arises in the radio band, inducing the QLD and subsequently creating relatively high pitch angles, which in turn lead to the X-ray emission.}\\label{fig1} \\end{figure} Most AGN exhibit VHE emission, which in turn indicates that the AGN magnetosphere is contained of ultra-relativistic electrons. In this context, the origin of acceleration of particles is very important. Proposed mechanisms such as Fermi-type acceleration (\\cite{cw99}), centrifugal acceleration (\\cite{mr94,osm7,ra08}), and acceleration due to the black hole dynamo mechanism (\\cite{levins}) can effectively provide Lorentz factors of the order of $\\sim 10^{5-9}$.  For an isotropic distribution of relativistic electrons (\\cite{blan}) in very strong magnetic fields, particles emit in the synchrotron regime with the power, $P\\approx 2e^4B^2\\beta_{\\perp}^2\\gamma^2/(3m^2c^3)$, therefore by assuming $\\beta_{\\perp}\\sim 1$, the synchrotron cooling timescale, $t_s\\equiv\\gamma mc^2/P$ can be estimated to be \\begin{equation}\\label{ts} t_s\\approx 5\\times 10^{-6}\\times\\left(\\frac{10^4G}{B}\\right)^2\\times\\left(\\frac{10^7}{\\gamma}\\right)s, \\end{equation} where we have taken into account that the energy is uniformly distributed between the beam and the plasma components, $n_b\\gamma_b\\approx n_p\\gamma_p$, \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics[angle=0]{13515fg2.eps}} \\caption{The synchrotron emission energy versus the cyclotron frequency. The set of parameters is $n_b = 2000cm^{-3}$, $L = 10^{45}erg/s$, $\\Omega = 3\\times 10^{-5}s^{-1}$, $\\gamma_p = 2$, $\\gamma_b\\in\\{4;5\\}\\times 10^6$ and $\\rho = r_{lc}$. As is clear, the low frequency cyclotron mechanism arises in the radio band, inducing via the QLD, the X-ray emission.}\\label{fig2} \\end{figure} and the magnetic field is normalized to the typical value of the magnetic induction close to the supermassive black hole (\\cite{paradig}). As is clear from this expression, the synchrotron timescale for ultra-relativistic electrons in the very vicinity of the black hole is so short ($\\sim 10^{-5}s$) that these particles very rapidly lose their perpendicular kinetic energy and transit to the ground Landau state. One can obtain a similar result for outer magnetospheric lengthscales, thus for the light cylinder area, when the magnetic field is determined by the bolometric luminosity, $L$ of AGN to be the equipartition field \\begin{equation} \\label{b} B_{lc}\\approx \\sqrt{\\frac{2L}{r_{lc}^2c}}\\approx 290\\times\\left(\\frac{L}{10^{45}erg/s}\\right)^{1/2}\\times\\left(\\frac{\\Omega}{3\\times 10^{-5}s^{-1}}\\right)G, \\end{equation} where $r_{lc} = c/\\Omega$ is the light cylinder radius and $\\Omega$ the angular velocity of rotation. If we consider a typical value of the bolometric luminosity, $\\sim 10^{45}erg/s$, then one can see from Eq. (\\ref{b}) that the magnetic induction equals approximately $290G$, which, combined with Eq. (\\ref{ts}), leads to the cooling timescale, $5.6\\times 10^{-4}s$.  All quantities vary because of the cyclotron instability, which causes to the QLD. By using Eqs. (\\ref{delta}-\\ref{om1}), we can estimate the frequency of the cyclotron mode to be \\begin{equation} \\label{om} \\omega\\approx 6.2\\times 10^8\\left(\\frac{\\gamma_p}{2}\\right)^4\\left(\\frac{10^7} {\\gamma_b}\\right)^2\\left(\\frac{B}{10^4G}\\right)^3\\left(\\frac{2\\times 10^3cm^{-3}}{n_b}\\right)Hz, \\end{equation} We consider a nearby zone of the AGN with mass $M_{BH} = 10^9M_{\\odot}$ ($M_{\\odot}$ is the solar mass) and typical magnetospheric parameters $n_b = 2000cm^{-3}$, $B = 10^4G$ and $\\gamma_p = 2$. Then, examining the ultra-relativistic beam component electrons with Lorentz factors $\\gamma_b\\in\\{3;4\\}\\times 10^7$ and assuming that the curvature radius $\\rho$ of field lines is of the order of the gravitational radius, $R_g\\equiv 2GM_{BH}/c^2$, one can see from Eq. (\\ref{pitch}) that the  created pitch angle is of the order of $10^{-9}rad$, which indicates the synchrotron emission in the $X$-ray domain [see Eq. (\\ref{eps})]. In Fig. \\ref{fig1}, we show the synchrotron emission energy versus the excited cyclotron frequency. As is clear from the figure, the cyclotron instability appears in the radio band leading to the $X$-ray emission. One can see that the radio emission is generated by the collective phenomena. We consider the radio frequency $5MeV$, corresponding to the wavelength, $\\lambda$, of the order of $6000 cm$. On the other hand, the average distance between particles, $d=n_p^{-1/3}$ is of the order of $2\\times 10^{-4}cm$, which is shorter by many orders of magnitude than $\\lambda$. This in turn, indicates that emission in the radio band is provided by the collective phenomena. In contrast to this, the VHE radiation generated in the $X$-ray domain (see Fig. \\ref{fig1}) is a single particle mechanism. The importance of the quasi-linear diffusion is twofold: (a) it generates the synchrotron radiation in strong magnetic fields for ultra-relativistic particles, which would be impossible without the QLD and (b) it simultaneously excites emission in two different domains. In this context, it is easy to check whether the emission is driven by the QLD by verifying that (I) the radiation is linearly polarized and (II) both signals have equal phases. In (\\cite{difus}), we performed the same theoretical analysis of the observed VHE ($>25GeV$) pulsed emission of the Crab pulsar (\\cite{magic}).  Unlike the previous case, we show in Fig. \\ref{fig2} the same behavior for the outer magnetospheric (light cylinder) lengthscales. The set of parameters is $M_{BH} = 10^9M_{\\odot}$, $n_b = 2000cm^{-3}$, $L = 10^{45}erg/s$, $\\Omega = 3\\times 10^{-5}s^{-1}$, $\\gamma_p = 3$, $\\gamma_b\\in\\{3;4\\}\\times 10^7$ and $\\rho = r_{lc}$. As shown in Fig. \\ref{fig2}, the radio frequency close to the light cylinder zone, in the kHz domain excites the hard $X$-ray emission by means of the QLD. From Eqs. (\\ref{g},\\ref{f}), one can straightforwardly check the validity of our assumptions, $|G_{\\perp}|\\gg |F_{\\perp}|$ and $|G_{_\\parallel}|\\ll |F_{_\\parallel}|$, confirming our approach. \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics[angle=0]{13515fg3.eps}} \\caption{The synchrotron emission energy versus the AGN luminosity. The set of parameters is $n_b = 2000cm^{-3}$, $L = 10^{45-47}erg/s$, $\\Omega = 3\\times 10^{-5}s^{-1}$, $\\gamma_p = 2$, $\\gamma_b = 5\\times 10^6$ and $\\rho = r_{lc}$. }\\label{fig3} \\end{figure} It is also interesting to investigate the QLD for different values of the luminosity. We examine the luminosity interval from $10^{45}erg/s$ to the Eddington limit, which for the given black hole mass, $10^9M_{\\odot}$, equals $10^{47}erg/s$. In Fig. \\ref{fig3}, we show the strong $X$-ray energy versus the AGN luminosity. The set of parameters is $M_{BH} = 10^9M_{\\odot}$, $n_b = 2000cm^{-3}$, $L = 10^{45-47}erg/s$, $\\Omega = 3\\times 10^{-5}s^{-1}$, $\\gamma_p = 3$, $\\gamma_b = 3\\times 10^7$ and $\\rho = r_{lc}$. As is evident from the figure, the emission energy is a continuously decreasing function of the luminosity. Indeed, by combining Eqs. (\\ref{eps},\\ref{A},\\ref{pitch},\\ref{b}), we see that the synchrotron emission energy behaves as $L^{-3/4}$.  Therefore, as our investigation shows, the QLD is a working mechanism in AGN magnetospheric flows and drives the synchrotron process.     The main aspects of the present work can be summarized as follows: \\begin{enumerate}   \\item We studied the quasi-linear interaction of proper modes of AGN magnetospheric plasmas with the resonant plasma particles. For this purpose, the synchrotron reaction force has been taken into account. The role of the QLD was studied in the context of producing the soft and hard $X$-ray emission from AGN.  \\item It has been shown that the synchrotron cooling timescales for ultra-relativistic electrons are very small, and particles rapidly transit to the ground Landau state, that in turn, prevents the subsequent radiation. We found that, under certain conditions the cyclotron instability develops, leading to the creation of pitch angles, and the subsequent synchrotron process.  \\item We have considered two extreme regions of magnetospheres: (a) relatively close to the black hole and (b) the light cylinder zone. As our model shows, the cyclotron instability, under certain conditions, generates the radio frequency in the range $(0.06-35)MHz$ and creates the soft $X$-ray emission, $(0.13-100)keV$, via the QLD. For the light cylinder area, radio spectra occurs in the range $(38-59)kHz$, which produces the hard $X$-ray emission in the domain $(1.8-2.9)keV$. We have emphasized that from an observational evidence one can directly verify the validity of the QLD by determining (I) the polarization and (II) phases of signals in radio and $X$-ray domains respectively.  \\item The quasi-linear diffusion has also been studied versus the AGN luminosity. It was shown that for more luminous AGN, the corresponding photon energy of the hard $X$-ray emission, generated by the synchrotron mechanism is lower.   \\end{enumerate}"
    },
    "0910/0910.3458_arXiv.txt": {
        "abstract": "We present the results of an Australia Telescope Compact Array 1.4\\,GHz spectropolarimetric aperture synthesis survey of 34 square degrees centred on Centaurus A---NGC~5128. A catalogue of 1005 extragalactic compact radio sources in the field to a continuum flux density of 3\\,mJy~beam$^{-1}$ is provided along with a table of Faraday rotation measures (RMs) and linear polarised intensities for the 28\\% of sources with high signal-to-noise in linear polarisation. We use the ensemble of 281 background polarised sources as line-of-sight probes of the structure of the giant radio lobes of Centaurus A. This is the first time such a method has been applied to radio galaxy lobes and we explain how it differs from the conventional methods that are often complicated by depth and beam depolarisation effects. Assuming a magnetic field strength in the lobes of $1.3\\,B_{1}\\,\\mu$G, where $B_{1}=1$ is implied by equipartition between magnetic fields and relativistic particles, the upper limit we derive on the maximum possible difference between the average RM of 121 sources behind Centaurus A and the average RM of the 160 sources along sightlines outside Centaurus A implies an upper limit on the volume-averaged thermal plasma density in the giant radio lobes of $\\langle n_e \\rangle < 5\\times10^{-5}B^{-1}_{1}$ cm$^{-3}$ . We use an RM structure function analysis and report the detection of a turbulent RM signal, with rms $\\sigma_{\\textrm{RM}} =17$\\,rad~m$^{-2}$ and scale size $0.3$\\degree, associated with the southern giant lobe. We cannot verify whether this signal arises from turbulent structure throughout the lobe or only in a thin skin (or sheath) around the edge, although we favour the latter. The RM signal is modelled as possibly arising from a thin skin with a thermal plasma density equivalent to the Centaurus intragroup medium density and a coherent magnetic field that reverses its sign on a spatial scale of 20\\,kpc. For a thermal density of $n_{1}\\,10^{-3}$ cm$^{-3}$, the skin magnetic field strength is $0.8\\,n^{-1}_{1}\\,\\mu$G. ",
        "introduction": "\\label{intro}     The lobes of radio galaxies are magnetised, quasi-freely expanding rarified cavities inflated by relativistic jets propagating outwards through the intergalactic medium from a central supermassive black hole (e.g. \\citealp{beg84}). As such, radio lobes could be excellent sites for high-energy particle acceleration and even the production of ultra-high-energy cosmic rays \\citep{ben08,fra08,hard09}. Knowledge of the physical conditions in radio lobes, including both the lobe magnetic field strength and thermal plasma density, are important to explore any high-energy acceleration mechanisms in full \\citep{kron04b}. Understanding the magnetic and thermal properties of radio lobes in detail is also fundamental to our understanding of galaxy formation in terms of feedback processes between the AGN and the interstellar/intergalactic medium \\citep{croton06,cro06,elbaz09}.\\newline  Linearly polarised electromagnetic radiation passing through a magnetised thermal plasma causes rotation in the angle of polarisation of the radiation at a rate given by, \\begin{equation}\\begin{split} \\delta\\theta &= 0.81~n_e~B_{\\parallel}~\\delta l~\\lambda^2\\\\ &= \\delta RM~\\lambda^2 \\label{eqn:rm1} \\end{split}\\end{equation} where $\\theta$ (in radians) is the position angle of the radiation at wavelength $\\lambda$ (in meters), $n_e$ is the thermal electron density (in cm$^{-3}$), $B_{\\parallel} $ is the line of sight component of the magnetic field (in $\\mu$G), $l$ is the path length through the rotating material (in pc) and RM is the Faraday rotation measure (in units of rad~m$^{-2}$). \\newline  If radio galaxy lobes contain magnetised, thermal plasma they will have an associated intrinsic Faraday depth. Observations of RMs in radio sources have shown that internal Faraday rotation in radio lobes is quite small \\citep{kron86,kron04b,sch98,palma00} with estimates of the thermal electron densities of $n_e\\lesssim 10^{-6}$ cm$^{-3}$, assuming a thermal plasma is distributed uniformly across the lobes. Such low inferred thermal matter densities are orders of magnitude lower than the upper limits on hot (keV) gas obtained from measurements of X-ray cavities around radio lobes \\citep{blanton01,fabian00,nulsen05}.  \\newline  Whereas only upper limits exist on the uniform thermal matter density inside radio lobes, there is conflicting evidence regarding the presence of a Faraday rotating thin skin (or sheath) around the lobes caused by entrainment of intergalactic plasma. For example, in the case of Cygnus~A, \\citet{dreh87} attribute observed RM variations of thousands of rad~m$^{-2}$ wholly to the foreground intracluster medium that Cygnus~A is embedded within. \\citet{bicknel90}, however, use the same data to show that this RM structure could arise from a thin skin around the radio lobes where Kevin-Helmholtz instabilities have caused a mixing of the lobe plasma with the intergalactic medium. More recently, a similar debate has arisen as to the origin (intracluster medium versus thin skin) of RMs (and RM variations) in excess of $\\pm$1000~rad~m$^{-2}$ across the radio galaxy PKS\\,1246$-$410 in the centre of the Centaurus cluster\\footnote{The Centaurus cluster at $z=0.01$ is located behind the Centaurus group at $z=0.0018$ that contains the radio galaxy Centaurus A.} \\citep{tay02,rud03,ens03,tay07a}. \\newline  There are three distinct scenarios we consider (but see \\citealt{burn66}) when using Faraday rotation to probe the properties of a magnetised, thermal plasma:\\newline  1. The diffuse polarised synchrotron emission from the lobes is mixed with the magnetised, thermal (Faraday rotating) plasma. In this case, the emission from the back of the source will have been rotated more than the emission from the front of the source and so `depth depolarisation' can occur along any line of sight. In addition, `beam depolarisation' can occur due to variations in RM on scales smaller than the observing beam.  In the former situation (i.e. not applicable to beam depolarisation) the concept of Faraday depth is introduced. Accurate determination of Faraday depth is complex, but necessary to extract information on the magnetic field and thermal density within the source \\citep{cioffi80}.  Radio galaxy lobes embedded in clusters or groups are often used to probe the cluster/group medium \\citep{dreh87,clarke01,eilek02,tay02,laing2008}. Here one must first show that the rotating plasma arises purely from the foreground medium itself rather than the lobes or skin of the radio source \\citep{rud03}.\\newline  2. The polarised emission is diffuse and located behind the Faraday rotating plasma. This is similar to the above scenario in terms of beam depolarisation, however no depth depolarisation occurs because there is no mixing of the emitting and rotating regions. For example, extended radio galaxies could be used to directly probe Galactic magnetic fields. \\newline  3. The polarised emission is unresolved and located behind the magnetised, thermal (Faraday rotating) plasma. In this scenario, all the polarised signal from any element of the background source is rotated by the entire line of sight through the screen. The screen can cause spatial depolarisation due to variations in RM in the screen across the angular size of the source. This probes scale sizes in the screen on scale sizes smaller than the background source size. No depth depolarisation occurs. This technique is often used to investigate the magnetic structure of the Milky Way \\citep{jb03,bro07,mao08}, nearby galaxies \\citep{han98,gaen05} and galaxy clusters \\citep{kim91,hennessy89}.\\newline  Note that scenarios 2 and 3 above are not affected by back-versus-front differences such as the Laing-Garrington effect \\citep{laing88,gar88}.\\newline  The typical angular size subtended by radio galaxy lobes is too small to include a statistically significant number of compact polarised background sources, at least for the source densities reached with current sensitivity (typically mJy~beam$^{-1}$). Hence, up until now, all studies of the Faraday rotation in radio galaxy lobes have been restricted to using emission from the lobes themselves, as in scenario 1 above (recent examples include \\citealp{kharb09,laing2008}). The outer lobes of the nearest radio galaxy, Centaurus A, subtend a large enough angular size ($\\approx 45$ deg$^2$) that hundreds of polarised sources are detected along sightlines behind them. For the first time, we can investigate the magnetised plasma in radio lobes --- using scenario 3 --- without the complexities added by depth and beam depolarisation effects. This is the basis for the analysis presented in this paper. \\newline   We have recently completed a large spectropolarimetric imaging campaign at 1.4\\,GHz with the Australia Telescope Compact Array (ATCA) and the Parkes 64\\,m radio telescope, to image in full the polarised structure of the nearest radio galaxy, Centaurus A. The full spectropolarimetric images of Centaurus A from ATCA and Parkes data combined will be reported in a subsequent paper (Feain et al. 2010, in preparation). An additional result of the ATCA component of the observations, we have also observed in full polarisation 1005 compact radio sources in the background of Centaurus A: some along lines of sight through the lobes and some along lines of sight beyond (outside) the boundaries of the radio lobes.  In this paper we present a catalogue of these 1005 compact radio sources. RMs and polarised flux densities are presented for a subset (281) of the sources. We investigate the spatial correspondence between the radio lobes of Centaurus A and the both the distributions of RMs and fractional polarisation of the background sources. \\newline  This paper is divided into sections in the following way: \\S\\ref{observations} describes the ATCA radio continuum observations and \\S\\ref{calibration} outlines the calibration and data processing procedures. In \\S\\ref{tara-catalogue}, we define the procedure used for source finding, give the format for the source catalogue and provide the URL where the entire catalogue can be accessed. \\S\\ref{polderivation} describes the Rotation Measure Synthesis technique with which we derived reliable RMs and polarised flux densities for 281 out of the 1005 sources catalogued. In \\S\\ref{src-counts} we show the RM distribution and briefly compare the total and polarised intensity source counts from our data with total and polarised intensity source counts from the literature. We also compute a lower limit on the very small scale RM fluctuations from the lobes of Centaurus A.  \\S\\ref{spatial-variation} presents a detailed investigation into the spatial variations in the ensemble of 281 Faraday rotation measures. We fit and subtract out the Milky Way foreground RM component leaving a residual excess RM dispersion on angular scales $\\theta\\sim0.3$\\degree. We model this excess as possibly arising from a thin skin around the southern giant radio lobe. Finally, our concluding remarks are given \\S\\ref{conclusion}.\\newline%  We adopt a distance to Centaurus A of 3.8\\,Mpc from \\citet{rej04}. At 3.8\\,Mpc, 1\\degree\\ corresponds to $\\approx$66\\,kpc. \\newline ",
        "conclusions": "\\label{conclusion}  In this paper, we have presented the results of a 1.4\\,GHz spectropolarimetric imaging survey of 34 deg$^2$ centred on the nearest radio galaxy Centaurus A. A catalogue of the 1005 background, compact radio sources brighter than 3\\,mJy~bm$^{-1}$ has been made available in electronic format. \\newline  We used Faraday rotation measure (RM) synthesis to measure linear polarised intensities and RMs for each source in our sample. There were 281 sources, out of 1005 in total, with a SNR$\\geq$7 in linear polarised intensity and for these we also publish a table of polarised intensities and RMs. The density of RMs published here, 8.3 per deg$^2$, represents the densest RM grid known to date.  Of the 281 RMs, 121 come from sources located behind the radio lobes of Centaurus A; the remaining 160 being from sources located along sightlines outside (in projection) the radio lobes. \\newline  The continuum source counts and the linear polarised intensity (P) source counts (for sources with SNR$\\geq$7 in P) are consistent with published source counts giving us confidence that our sample is statistically robust and the emission from the lobes of Centaurus A has not contaminated the global statistics of our sample. There is no measurable difference between the polarised source counts behind the lobes and those outside the lobes. We derived an upper limit on the very small scale turbulence in the lobes of $<10$\\,rad\\,m$^{-2}$ on scales $\\lesssim$180\\,pc. This is similar to the limit derived on the volume-averaged RM signal from the lobes. \\newline  We modelled the Milky Way contribution to the RMs by fitting a first order polynomial surface to the 160 RMs along sightlines outside the lobes of Centaurus A and subtracting this surface from all 281 RMs in our sample. The residual RMs behind the lobes have a mean of $-4.5\\pm2.8$~rad\\,m$^{-2}$ and a standard deviation of 30.4~rad\\,m$^{-2}$, compared to a mean of $-1.6\\pm2.1$~rad~m$^{-2}$ and standard deviation of 26.6~rad\\,m$^{-2}$ for RMs outside the lobes. Assuming a magnetic field strength in the lobes of $1.3\\,B_{1}\\,\\mu$G, we derived an upper limit to the uniform thermal plasma density in the lobes of $\\langle n_e \\rangle < 5\\times10^{-5}\\,B_1$ cm$^{-3}$. This is a factor of $2-3$ times smaller than previously published limits and further constrains the prospects for high-energy acceleration processes in the lobes of Centaurus A.\\newline  We used a structure function analysis to measure an excess RM dispersion in the southern lobe of Centaurus A of $\\sigma_{\\textrm RM}\\approx 17~$rad\\,m$^{-2}$ on angular scales of $\\sim0.3$\\degree\\ associated with the southern lobe of Centuarus A. The RM signal is modelled as possibly arising from a thin skin with a thermal plasma density equivalent to the Centaurus intragroup medium density and a coherent magnetic field that reverses its sign on a spatial scale of 20\\,kpc. If the Centuarus intragroup density is of order $10^{-3}$ cm$^{-3}$, the skin magnetic field strength is of order $1~\\mu$G. If $n_e$ is an order of magnitude smaller, $B_t$ will be $10\\times$ larger. \\newline  A more sensitive survey of the southern lobe is now required to verify if the RM signal is associated with the entire southern lobe, or a thin skin (sheath) around the edges, although the latter seems more likely.  In addition, a deep X-ray observation of the edge of the southern lobe would be incredibly revealing in terms of the detailed physics it could probe. The southern lobe of Centaurus A has typically been neglected in favour of investigations of the northern lobe; this is almost certainly because of the very interesting northern middle lobe (NML) region (there has been no southern counterpart to the NML detected). In this paper, we have shown that the edge of the giant southern lobe, rather than any part of the giant northern lobe, is probably associated with fluctuating RMs on scales of 20\\,kpc. It is extremely timely now to follow-up this region with a multiwavelength campaign at, in particular, X-ray and deep narrow-band optical wavelength regimes. We are currently completing a detailed analysis of the full polarisation radio continuum images of Centaurus at a resolution of $\\approx 40$\\arcsec. In this analysis, we will be able to compare the fluctuating RM region with the polarised and total intensity structure of the lobes at a factor of about 5 times the resolution that has been previously possible. If our model of a thin skin is correct, these new images should reveal the signatures of surface waves around the lobes, which would be evidence of mixing between the lobe emission and the Centaurus intragroup medium.\\newline"
    },
    "0910/0910.4108_arXiv.txt": {
        "abstract": "The broad Balmer emission-line profiles resulting from the reprocessing of UV/X-ray radiation from a warped accretion disc induced by the Bardeen-Petterson effect are studied. We adopt a thin warped disc geometry and a central ring-like illuminating source in our model. We compute the steady-state shape of the warped disc numerically, and then use it in the calculation of the line profile. We find that, from the outer radius to the inner radius of the disc, the warp is twisted by an angle $\\sim\\pi$ before being flattened efficiently into the equatorial plane. The profiles obtained depend weakly on the illuminating source radius in the range from $3r_{g}$ to $10r_g$, but depend strongly on this radius when it approaches the marginally stable orbit of an extreme Kerr black hole. Double- or triplet-peaked line profiles are present in most cases when the illuminating source radius is low. The triplet-peaked line profiles observed from the Sloan Digital Sky Survey may be a {``}signature\" of a warped disc. ",
        "introduction": "\\label{intro}  A small fraction of active galactic nuclei~(AGNs) with broad lines show double-peaked emission-line profiles, which suggests an accretion disc origin \\citep*{era94,era03,str03,gez07}. In a stationary circular relativistic disc model, Doppler boosting will make the blue peak of the profile higher than the red one. However, observations of the variability of line emission in some objects show that at least at some epochs the profile asymmetry is reversed, with the red peak higher than the blue one, contrary to the predictions of homogeneous, circular relativistic disc models \\citep{mil90,str03,gez07}. To adjust the theoretical to the observed profiles, non-axisymmetric or warped discs are usually required. Various physically plausible processes in the accretion flow have been considered, such as spiral shocks, eccentric discs, bipolar outflows or a binary black hole system \\citep{cha94,era95,sto97,zhe90,vei91,beg80, gas96,zha07b}, as well as a warped accretion disc \\citep*{wu08}. This warped disc makes it possible for the radiation from the inner disc to reach the outer parts, enhancing the reprocessing emission lines. Thus, warping can both provide the asymmetry required for the variations of the emission lines and naturally solve the long-standing energy-budget problem, because the subtending angle of the outer disc portion to the inner one is increased by the warp.  There are strong observational and theoretical grounds for believing that accretion discs around black holes may be warped and twisted. Evidence for the existence of warped discs in astrophysical systems has been found from observations \\citep{her05,cap06,cap07,mar08}. From a theoretical point of view, four main mechanisms for exciting/maintaining warping in accretion discs have been proposed, namely warping that is tidally induced by a companion in a binary system \\citep{teb93,teb96,lar96}, radiation-driven or self-induced warping \\citep{pri96,pri97,mal96,mal97,mal98}, magnetically driven warping \\citep{lai99,lai03,pfl04}, or warping driven by frame dragging \\citep{bap75,kum85,arm99}. Note also that warps generated by gravitational interactions have been investigated in the literature. In the galactic context, \\citet{hun69} studied the linear bending waves of a self-gravitating, isolated, thin disc. On nuclear disc scales, \\citet{pap98} studied the evolution of a thin self-gravitating viscous disc interacting with a massive object orbiting the central black hole. \\citet{ulu09} showed that highly warped discs near black holes can persist for a long time without any persistent forcing other than by their self-gravity. As a result, the accretion disc in some AGNs may be non-planar.  To solve the energy budget for the emission-line region for the disc model, reprocessing of the ultraviolet(UV)/X-ray continuum from the inner accretion disc is required. However, only a very small fraction of radiation from the inner disc is expected to intercept the outer part of the disc in the case of a flat geometrically thin disc. Two types of processes have been proposed to increase the fraction of light incident on the disc emission-line region. \\citet*{che89} proposed that the inner part of the accretion disc is hot and becomes geometrically thick as a result of insufficient radiative cooling, and the X-ray emission from such an ion-supported torus is responsible for such energy input. However, double-peak emitters do not always have a low accretion rate \\citep{zha07a,bian07}. In contrast, \\citet{cao06} proposed that a slow-moving jet can scatter a substantial fraction of UV/X-ray light from the inner accretion disc back to the outer part of the disc. Although this may be likely for radio-loud double-peaked emitters, the majority of double-peaked emitters are radio-quiet. In this work, we consider a system with an accretion disc around a Kerr black hole. The black hole spin is misaligned with the outer parts of the disc, as there is no physical reason to suppose that the angular momentum of the accreting mass and the angular momentum of the spinning black hole are always aligned. Frame dragging produced by a Kerr black hole causes the precession of the orbit of a particle if its orbital plane is inclined with respect to the equatorial plane of the black hole. This is known as the Lense-Thirring effect. The combined action of the Lense-Thirring effect, which tends to twist up the disc, and the internal viscosity of the accretion disc, which tends to smooth it out, forces the alignment between the angular momenta of the Kerr black hole and the accretion disc. This is known as the Bardeen-Petterson effect. This effect tends to affect the inner regions of the disc owing to the short range of the Lense-Thirring effect($\\propto r^{-3}$), whereas the outer part of the disc tends to keep its original inclination. The transition radius between these two regimes is known as the Bardeen-Petterson radius $R_{\\rm BP}$~(or warp radius $R_{\\rm w}$), and its exact location depends mainly on the physical properties of the accretion disc \\citep{bap75,sch96,nel00,fra05a}. Basically, the Bardeen-Petterson radius is determined by comparing the time-scale related to the Lense-Thirring precession with that of warp propagation through the disc. We consider here the propagation of warps in thin Keplerian discs in the case for which the disc is sufficiently viscous that the warp propagates in a diffusive manner. The twisted disc itself is treated using a non-relativistic approach with an additional term describing the gravitomagnetic precession of the disc rings. Both analytic and numerical steady-state solutions have shown that the evolution of a misaligned disc arising from the Bardeen-Petterson effect usually produces an inner flat disc and a warped transition region with a smooth gradient in the tilt and twist angles \\citep{sch96,nat99,mar07}. If the disc is thick and/or its viscosity is low, the warp propagates in a wave-like rather than a diffusive manner \\citep*{nel00,lub02}. In this case, the twisted disc is not necessarily aligned with the equatorial plane at small scales. The inclination angle of a low-viscosity disc may even oscillate or counteralign close to the black hole \\citep{lub02,kin05,lod06,lod07}.  The effect of a Bardeen-Petterson disc on iron line profiles treated relativistically has been examined by \\citet{fra05b}, and \\citet{bac99} used a non-relativistic treatment to study the broad-line H$\\beta$ profiles from a warped disc. A relativistic treatment of the effect of a warped disc with a parametrized geometry on the Balmer lines that can be compared with observations was presented by \\citet{wu08}. Following on from these calculations, the disc warping induced by the Bardeen-Petterson effect and its influence on the broad Balmer emission owing to the reprocessing of the central high-energy radiation is investigated. In \\s~\\ref{method} we summarize the assumptions behind our model and present the basic equations relevant to our problem. We present our results in \\s~\\ref{result}, and the conclusions and a discussion in \\s~\\ref{sum}. ",
        "conclusions": "\\label{sum} We computed the Balmer emission-line profiles that result from the reprocessing of emission using a numerical model of a warped disc induced by the Bardeen-Petterson effect, including all relativistic effects for radiation propagation, and subject to various shadowing effects associated with disc warping. We numerically solved the disc structure for a steady-state disc shape. To simplify the analysis, we took both viscosities to have the same power-law form, and thus the ratio $\\nu_2/\\nu_1$ is independent of radius. For simplicity, we assumed that the disc is illuminated by a ring-like central source, which is a rough approximation, that the line emissivity is proportional to the continuum light intercepted by the accretion disc, and that line emission is isotropic. Our conclusions are as follows. \\begin{enumerate} \\item From the numerical results, we can see that, moving inwards from the outer radius of the disc, the warp is twisted by an angle of $\\sim\\pi$ before being flattened efficiently into the equatorial plane. \\item For a given warped disc, there are two angles that determine the observed line profile, the inclination angle and the azimuthal viewing angle resulting from the non-axisymmetry of the disc warping. \\item For less twisted warped disc induced by the Bardeen-Petterson effect, the asymmetrical and frequency-shifted single-peaked line profiles are produced in most cases when the primary source radius is large(from 3$r_g$ to 10$r_g$). The double- or triplet-peaked line profiles presented in most cases occur when the primary source radius approaches the marginally stable orbit of an extreme Kerr black hole.  \\end{enumerate}  The line profiles depend strongly on the twisting structure of the disc \\citep{wu08} as well as on the radius of the illuminating source. The Doppler effect and gravitational focusing concentrate the radiation from the innermost region towards the equatorial plane, which increases the returning flux \\citep{cun75,cun76}. This may be the reason why the illuminating source radius has a strong influence on the line profile. Knowledge of the twisting structure may allow the determination of some important characteristics of AGNs, such as the spin of the black hole and the viscosities of the disc. We assume that the ratio $\\nu_2/\\nu_1$ is independent of radius in the calculations, but the relation between $\\nu_1$ and $\\nu_2$ is very uncertain, and may influence the disc twisting structure. It is also important to emphasize that, although we have analyzed the Bardeen-Petterson effect in this work, the four mechanisms mentioned above are not mutually exclusive. This means that more than one mechanism might be operating in a given system."
    },
    "0910/0910.0806_arXiv.txt": {
        "abstract": " ",
        "introduction": "One of the goals of modern cosmology is to understand and quantify the nature of large scale matter distribution of the Universe. An accurate empirical description of large scale clustering of matter, derived from systematic observations of visible matter in the Universe, is   essential to achieve this goal. Efforts in this directions have vastly improved due to better instruments that have been available in recent times for astronomical data acquisition as well as better statistical techniques of analysis of the acquired data. It is as a result of these efforts  that an enormous amount of data about the observable universe has been accumulated in the form of the now well-known {\\it redshift surveys}, and some widely accepted conclusions  drawn from these data have created a certain confidence in many researchers that an accurate description of the large scale matter distribution is just about being achieved.  In statistical analysis of galaxy distribution, we are not interested in the number of galaxies in a particular region of the sky but we are rather interested only in the average properties of  number distribution of galaxies. We are {\\it e.g.} interested in knowing whether or not distribution of galaxies is clumpy, and if so, how we can quantify the nature of clumpiness. In the literature \\citep{2002sgd..book.....M}, various statistical methods of analysis of galaxy clustering have been discussed. The broad feature of all these methods is to discuss the nature of clumpiness of the galaxy distribution. Among all these methods the historical favorites have been variants of two point correlation function (see equation \\ref{eq21}). This function measures the excess probability, relative to a Poisson distribution, of finding an object near another object. Bok's statistic (the dispersion of the counts $N$ in cells), is an integral over two point correlation function. Zwicky's index of clumpiness is the ratio of variance of $N$ to what would be expected for a uniform random distribution. From the two point correlation function of the counts of galaxies in the Lick survey, Limber showed that there is a linear integral equation relating the angular correlation function to the corresponding spatial correlation function. Neyman and Scott devised a {\\it priori} statistical model of clustering and then adjusted the parameters to fit model statistics to estimates from data. A recent program in a similar vein is called the halo model. Recently there have been precise estimates of two point correlation function from redshift surveys like Two Degree Field Galaxy Redshift Survey (2dFGRS) and the Sloan Digital Sky Survey (SDSS).  The Fourier or spherical harmonic transform of the two point correlation function is the power spectrum (see equation \\ref{eq221}). It is the description of clustering in terms of wavenumbers $k$ that separates the effects of different scales. Other descriptors of statistics have been the $N$ point correlation function of the distribution of points, moments and counts in cells, void probability function and nearest neighbor  distances etc. However all these methods are based on the idea of eventual homogenization of the matter distribution within the sample size itself. A group of researcher feels that this idea of homogenization is flawed and the distribution of matter in the Universe is intrinsically inhomogeneous to largest observed scales and, perhaps, indefinitively.  The debate of homogeneous versus non homogeneous distribution of matter has taken a new vigor with the arrival of a new method for describing the clustering of galaxies. This method is based on the ideas of a new geometrical perspective for the description of irregular patterns in nature. We generally refer to it as \\textit{the fractal geometry}. In this chapter we intend to show the basic idea behind this geometrical approach.  The plan of this chapter is as follows. In section \\ref{Standard Analysis} we  briefly present the basic tools used in statistical analysis of the large scale distribution of galaxies, its estimations, difficulties and answers given to these difficulties. The subsection \\ref{Higher Order} describes Higher order statistics of the distribution. Section \\ref{other statistical} describes various other methods of statistical analysis. Section \\ref{Fractals}  presents a brief, but general, introduction to fractals, which emphasizes their empirical side and applications. The discussion on various fractal dimensions along with their merits (and demerits) follow in section \\ref{Other Frac Dim} and \\ref{Prob}. We conclude the chapter with a small discussion on Lacunarity of the point distribution in section \\ref{lacunarity}. ",
        "conclusions": "We have studied the problem of the expected value of the Minkowski-Bouligand dimension for a finite distribution of points. For this purpose, we have studied a homogeneous distribution as well as a weakly clustered distribution. In our study, $q/\\bar{N}$ and $q\\bar\\xi$ are taken to be small parameters and the deviation of $D_q$ from $D$ is estimated in terms of these quantities. In both cases we find that the expected values of the Minkowski-Bouligand dimension $D_q$ are different from  $D$ for the distribution of points. For generic distributions, the value of $D_q$ is less than the dimension $D$. We have derived an expression for $D_q$ in terms of the correlation function and the number density in the limit of weak clustering. We see that the dimension of the ambient space (i.e. $D$) is not the correct benchmark for defining a homogeneous distribution and instead the minkowski bouligand dimension of the distribution should match with equation \\ref{eqn:homogen} for the distribution to be homogeneous.  We find that $\\Delta D_q = D_q - D$ is non-zero at all scales for unbiased tracers of mass in the concordance model in cosmology. For this model $\\Delta D_q$ is a large negative number at small scales but it rapidly approaches zero at larger scales. $\\Delta D_q$ is a very slowly varying function of scale above $100$~h$^{-1}$Mpc and hence this may be tentatively identified as the scale of homogeneity for this model. A more quantitative approach requires us to estimate not only the systematic offset $\\Delta D_q$ but the dispersion in this quantity. The scale of homogeneity can then be identified as the scale where the offset is smaller than the expected dispersion.  Although we have used the example of galaxy clustering for illustrating our calculations, the results as given in Equation (\\ref{eqn:homogen}) and Equation (\\ref{eqn:clus}) are completely general and apply to any distribution of points with weak departures from homogeneity. A detailed derivation of the relations presented here, with verification using mock distributions of points will be presented in a future publication, where we also expect to highlight other applications.  We have obtained the scale of homogeneity of the galaxy distribution in the volume limited samples obtained from first data release of Sloan Digital Sky Survey. We will discuss this in detail in chapter \\ref{chap5}. \\chapter{Testing homogeneity on large scales in the Sloan Digital Sky Survey}\\label{chap5} One of the most important assumption on which modern cosmology is based is the Cosmological Principle. It states that the Universe is homogeneous and isotropic on large scales. We\\footnote{This chapter is based on {\\it Testing homogeneity on large scales in Sloan Digital Sky Survey Data Release One}, (Jaswant Yadav, S. Bharadwaj, B. Pandey \\& T.R. Seshadri), Monthly Notices of Royal Astronomical Society, 2005, {\\bf 364}, 601} have tested the large scale homogeneity of the galaxy distribution in the Sloan Digital Sky Survey Data Release One (SDSS-DR1) using volume limited subsamples extracted from the two equatorial strips. These strips are nearly two dimensional (2D). The galaxy distribution is projected  on the equatorial plane and we have carried  out a 2-dimensional  multi-fractal analysis  by counting the number of galaxies inside circles of different radii $r$ in the range $5 \\, h^{-1} {\\rm Mpc}$ to $150 \\, h^{-1} {\\rm Mpc}$. The circles  have been centred on galaxies. Different moments of the count-in-cells have been analyzed to identify a range of length-scales ($60-70 \\, h^{-1} {\\rm Mpc}$ to $150 h^{-1} {\\rm Mpc}$ ) where the moments show a power law scaling behavior.  It has helped us to determine the scaling exponent which gives the spectrum of generalized dimension $D_q$. If the galaxy distribution is homogeneous, $D_q$ does not vary with $q$ and is equal to the Euclidean dimension which in our case is 2. We find that $D_q$ varies in the range $1.7$ to $2.2$. We also constructed mock data from random, homogeneous point distributions and from \\lcdm~ N-body simulations with bias $b=1, 1.6$ and $2$,  and analyzed these in exactly the same way. The values of $D_q$ in the random distribution and the unbiased simulations show much smaller variations  and these are not consistent with the actual data.  The biased simulations, however, show larger variations in $D_q$ and these are  consistent with both the random and the actual data. Interpreting the actual data as a realization of a biased \\lcdm~ universe, we conclude that the galaxy distribution is homogeneous on scales larger than  $60-70 \\, h^{-1} {\\rm Mpc}$."
    },
    "0910/0910.0171_arXiv.txt": {
        "abstract": "The body-centered cubic Coulomb crystal of ions in the presence of a uniform magnetic field is studied using the rigid electron background approximation. The phonon mode spectra are calculated for a wide range of magnetic field strengths and for several orientations of the field in the crystal. The phonon spectra are used to calculate the phonon contribution to the crystal energy, entropy, specific heat, Debye-Waller factor of ions, and the rms ion displacements from the lattice nodes for a broad range of densities, temperatures, chemical compositions, and magnetic fields. Strong magnetic field dramatically alters the properties of quantum crystals. The phonon specific heat increases by many orders of magnitude. The ion displacements from their equilibrium positions become strongly anisotropic. The results can be relevant for dusty plasmas, ion plasmas in Penning traps, and especially for the crust of magnetars (neutron stars with superstrong magnetic fields $B \\gtrsim 10^{14}$~G). The effect of the magnetic field on ion displacements in a strongly magnetized neutron star crust can suppress the nuclear reaction rates and make them extremely sensitive to the magnetic field direction. ",
        "introduction": "\\label{introduct} The model of a crystal of point-like charges immersed in a uniform neutralizing background of opposite charge was conceived by Wigner \\cite{W34} to describe a possible crystallization of electrons. These Wigner crystals of electrons have much in common with Coulomb crystals of ions with the uniform electron background. The model of the Coulomb crystal is widely used in different branches of physics, including theory of plasma oscillations (e.g., \\cite{P65}), solid state physics, and works on dusty plasmas and ion plasmas in Penning traps (e.g., \\cite{IBTJHW98}).  Moreover, Coulomb crystals of ions, immersed in an almost uniform electron background, are formed in the cores of white dwarfs and envelopes of neutron stars. The properties of such crystals are important for the structure and evolution of these astrophysical objects (e.g., \\cite{CADW92,HPY07}). In particular, the Coulomb crystal heat capacity  \\cite{BPY01} controls cooling of old white dwarfs and is used to determine their ages (e.g., \\cite{WKCM09}). Crystallization of white dwarfs can influence their pulsation frequencies. It can thus be studied by powerful methods of asteroseismology \\cite{MMK04}. Microscopic properties of Coulomb crystals determine the efficiency of electron-phonon scattering in white dwarfs and neutron stars, and, hence, transport properties of their matter (such as electron thermal and electric conductivities, and shear viscosity, e.g., \\cite{BKPY98,CY05}) as well as the neutrino emission in the electron-ion bremsstrahlung process \\cite{KPPTY99}. Many-body ion correlations in dense matter produce screening of ion-ion (nucleus-nucleus) Coulomb interaction and affect nuclear reaction rates in the thermonuclear burning regime with strong plasma screening and in the pycnonuclear burning regime (when the reacting nuclei penetrate through the Coulomb barrier owing to zero-point vibrations in crystalline lattice). The description of various nuclear burning regimes and observational manifestations of burning in white dwarfs and neutron stars  are discussed in Refs.\\ \\cite{SVH69,YGABW06,CD09} and references therein. The manifestations include type Ia supernova explosions of massive accreting white dwarfs, bursts and superbursts, deep crustal heating of accreted matter in neutron stars.  Since the late 1990s certain astrophysical applications have been requiring a comprehensive study of Coulomb crystals in strong magnetic fields. The topic has become important due to the growing observational evidence that some very intriguing astrophysical objects, soft-gamma repeaters (SGRs) and anomalous X-ray pulsars, belong to the same class of sources called magnetars (see, e.g., \\cite{T02,K04,WT06} and reference therein). These are thought to be isolated, sufficiently warm neutron stars with extremely strong magnetic fields $B \\gtrsim 10^{14}$~G. For instance, the magnetic field of SGR 1806--20, inferred from measurements of its spin-down rate, is $B \\sim 2 \\times 10^{15}$ G \\cite{catalog}. Magnetars are observed in all ranges of the electromagnetic spectrum. They show powerful quasipersistent X-ray emission, bursts and giant bursts with enormous energy release. During giant bursts, one often observes quasi-periodic X-ray oscillations which are interpreted (e.g., \\cite{WS07}) as vibrations of neutron stars (involving torsion vibrations of crystalline neutron star crust). It is likely that the activity of magnetars is powered by their superstrong magnetic fields. Thus the magnetars can be viewed as natural laboratories of dense matter in magnetic fields. In order to build adequate models of magnetar envelopes and interpret numerous observations, it is crucial to know the properties of magnetized Coulomb crystals. The main goal of this paper is to study in detail Coulomb crystals in an external uniform magnetic field. (The results reported here were partially presented in \\cite{B00}.)  The Coulomb crystals in question consist of fully ionized ions with charge $Ze$ and mass $M$ arranged in a crystal lattice and immersed into the rigid electron background (in this case, rigid means unpolarizable or incompressible, i.e.\\ constant and uniform). The effect of the magnetic field $B$ on the ion motion can be characterized by the ratio \\begin{equation} b = \\omega_B / \\omega_{p}, \\label{rhoB} \\end{equation} where \\begin{equation} \\omega_B=\\frac{ZeB}{Mc}, \\quad \\omega_{p}=\\sqrt{\\frac{4 \\pi Z^2 e^2 n}{M}} \\end{equation} are the ion cyclotron frequency and the ion plasma frequency, respectively; $n$ is the ion number density, while $c$ is the speed of light. It is expected that the magnetic field modifies the properties of the ion crystal at $b \\gtrsim 1$ (see, however, Figs.\\ \\ref{b-c} and \\ref{u1um1} below). In a strong magnetic field the approximation of rigid electron background is a bigger idealization of the real situation in neutron star crust matter, since higher densities are required to achieve full ionization and suppress the polarizability of the electron background. The higher densities (and $\\omega_{p}$) imply smaller $b$. However, the effective ion charge approximation turns out to be successful for analyzing partially ionized systems (e.g., \\cite{PCY97}). The quantity $b$ (also equal to $v_{A}/c$, where $v_{A} = B/\\sqrt{4 \\pi \\rho}$ is the Alfv\\'en velocity, $\\rho$ being the mass density) is, actually, independent of the ion charge. Hence, one can consider large $b$ in Coulomb crystals at not too high densities having in mind the effective ion charge approximation. Despite that, the effect of the polarizability of the compensating electron background has to be studied separately.  The closely related problem of magnetized Wigner crystals of electrons was studied in the early 1980s by Usov, Grebenschikov and Ulinich \\cite{UGU80} and by Nagai and Fukuyama \\cite{NF82, NF83}. Usov et al.\\ \\cite{UGU80} obtained the equations for crystal oscillation modes, studied qualitatively the oscillation spectrum, and diagonalized the Hamiltonian of the crystal for a proper quantum description of the oscillations in terms of phonons. In addition, the authors investigated asymptotic temperature and magnetic field dependences of the specific heat, rms electron displacement from the respetive lattice site, and magnetic moment of the crystal. They obtained a soft phonon mode with the frequency $\\Omega \\propto k^2$ near the center of the Brillouin zone. It resulted in a rather unusual low-temperature specific heat behavior ($T^{3/2}$) instead of the standard Debye law ($T^3$). Also, the authors mentioned the dependence of the crystal energy on the magnetic field direction as well as the increased stability of the crystal due to a suppression of electron vibration amplitudes. In addition, Usov et al.\\ \\cite{UGU80} considered the dielectric function of the crystal and studied the effect of the electromagnetic field induced by the electron motion. Their main results were, however, of semi-qualitative nature, limited to various extreme cases, whereas the present paper focuses mostly on quantitative results, pertaining to Coulomb crystals of ions, with an eye to astrophysical implications.  Nagai and Fukuyama \\cite{NF82} calculated phonon spectra of magnetized body-centered cubic and face-centered cubic (bcc and fcc) Wigner crystals and compared the energies of these crystals at zero temperature as a function of the magnetic field. The crystal energy was calculated as a sum of the electrostatic (Madelung) energy, and the zero-point energy of crystal vibrations. The effect of the anharmonic and exchange terms was neglected. The electrostatic energy is independent of the magnetic field (as long as the field does not alter the lattice structure). The vibration energy does depend on the field because the field modifies phonon modes. At zero field the energy minimum is realized by the bcc structure, both with and without zero-point vibrations (in general, it is not true for polarizable background, e.g., \\cite{B02}). The authors showed that for a sufficiently strong magnetic field and at relatively high densities [$r_s  < 275$, where $r_s = a_e/a_0$ is the standard density parameter, $a_e=(3/4 \\pi n_e)^{1/3}$, and $a_0$ is the Bohr radius, $n_e$ being the electron number density] the full energy is minimized by the fcc structure. It is worth mentioning that the authors did not consider the dependence of the zero-point energy on the magnetic field direction and performed all calculations for a fixed direction (along one of the high symmetry axes of the crystals). However, this choice of the field direction for the bcc lattice was not optimum, as far as the lattice energy was concerned. Moreover, the energy gain obtained by choosing the optimum direction is of the same order of magnitude as the difference between the zero-point energies of bcc and fcc lattices.  Nagai and Fukuyama \\cite{NF83} investigated an analogous structural transition between bcc and hexagonal close-packed (hcp) electron Wigner lattices. The hcp lattice was found energetically favorable for a sufficiently strong magnetic field at $r_s < 10700$. This result seems more robust since the energy difference between bcc and hcp lattices is several times larger than the energy gain obtained by choosing the optimum direction of the magnetic field in the bcc lattice (the field direction adopted in \\cite{NF83} was the same as in \\cite{NF82}).  In addition, Nagai and Fukuyama \\cite{NF83} described the behavior of all 6 phonon modes of the hcp lattice in the magnetic field. Also, they analyzed quantitatively the behavior of transverse and longitudinal electron displacements from the equilibrium positions for bcc and hcp lattices (at zero temperature) and concluded (qualitatively) that in strong magnetic fields the crystals became significantly more stable in the transverse direction and somewhat more stable in the longitudinal direction.  The present paper is organized as follows. Section \\ref{eqs} discusses equations for Coulomb crystal oscillations in the magnetic field. Section \\ref{spectrum} focuses on the phonon spectrum of such a crystal with a simple lattice (i.e., one ion in the primitive cell). The Hamiltonian of the system is diagonalized in Sec.\\ \\ref{diag}, which allows one to express the ion displacement operator via phonon creation and annihilation operators. Section \\ref{thermodyn} presents numerical calculations of thermodynamic functions of the Coulomb bcc crystal for a wide range of densities, temperatures and magnetic fields. In Sec.\\ \\ref{crust_ex} these results are applied to the real physical system found in magnetar crust. Section \\ref{rms} is devoted to an analysis of the rms amplitudes of ion vibrations and to the Debye-Waller factor of the magnetized Coulomb crystal. Finally, Sec.\\ \\ref{moments} considers the dependence of phonon spectrum moments on the magnitude and direction of the magnetic field.  The results will be parameterized by the quantum parameter $\\theta=\\hbar \\omega_{p}/T \\equiv T_{p}/T$ (where $T$ and $T_{p}$ are temperature and ion plasma temperature, respectively), and by the Coulomb coupling parameter $\\Gamma = Z^2 e^2 /(aT)$. In this case, $a=(3/4 \\pi n)^{1/3}$ is the ion sphere radius. The Boltzmann constant is set equal to $k_{B}=1$. ",
        "conclusions": "\\label{concl} The Coulomb crystal of ions with incompressible charge compensating background of electrons in constant uniform magnetic field has been studied. The phonon mode spectrum of the crystal with bcc lattice has been calculated for a wide range of magnetic field strengths and orientations (Fig.\\ \\ref{spectinb}). The phonon spectrum has been used in 3D numerical integrations over the first Brillouin zone for a detailed quantitative analysis of the phonon contribution to the crystal thermodynamic functions, Debye-Waller factor of ions, and the rms ion displacements from the lattice nodes for a broad range of densities, temperatures, chemical compositions, and magnetic fields.  The characteristic parameter that determines the strength of the magnetic field effects in the crystal is the ratio of the ion cyclotron frequency to the ion plasma frequency, $b = \\omega_B/\\omega_{p}$. Even a very strong field ($b \\gg 1$) does not alter the partition function (hence, thermodynamic functions, melting etc.) of a classical crystal, i.e.\\ a crystal at a temperature significantly exceeding the energy of any phonon mode ($T \\gg \\hbar \\omega_B$ for strongly magnetized crystal).  Strong magnetic field dramatically changes various properties of quantum crystals, especially at $\\theta = \\hbar \\omega_{p}/T \\gg 1$. Low-temperature phonon specific heat, entropy, thermal energy increase by orders of magnitude (e.g., Figs.\\ \\ref{b-c}, \\ref{b-se}). The thermodynamic functions exhibit peculiar staircase structures, brought about by the vastly different energy scales of the crystal phonon mode. Ion displacements from the equilibrium positions become strongly anisotropic (Fig.\\ \\ref{dw}), and so does the Debye-Waller factor.  These results can be applied directly to real physical systems found in the crust of magnetars (neutron stars with superstrong magnetic field). In particular, the heat capacity of the magnetized Coulomb crystal composed of fully ionized $^{56}$Fe has been analyzed (Sec.\\ \\ref{crust_ex}). It has been shown that in the magnetic field (as in the field-free case) the phonon heat capacity dominates the electron one in a wide range of densities and temperatures (Fig.\\ \\ref{iron_spec_heat}). In accordance with Sec.\\ \\ref{thermodyn}, magnetic field is responsible for a significant boost of phonon heat capacity, which may be important for simulations of magnetar cooling. Since ion displacements are suppressed, one can expect increased crystal stability in the magnetar crust associated with an increase of the melting temperature in the quantum regime. Finally, one expects a reduction of the pycnonuclear reaction rates due to increased tunnelling lengths. The rates are likely to become very sensitive to the orientation of the magnetic field (Sec. \\ref{rms}).  The dependence of thermodynamic functions on the magnetic field orientation within the crystal turns out to be weak, at least for the bcc lattice, but finite. The energy gain, achieved by reorienting the crystal with respect to the magnetic field, can be estimated using Fig.\\ \\ref{u1um1}.  The techniques used in the present paper can be easily reformulated for fcc lattice, where results are expected to be numerically close to those for bcc lattice."
    },
    "0910/0910.0492_arXiv.txt": {
        "abstract": "We use the direct Fourier method to calculate the redshift-space power spectrum of the maxBCG cluster catalog \\citep{2007astro.ph..1265K} -- currently by far the largest existing galaxy cluster sample. The total number of clusters used in our analysis is $12,616$. After accounting for the radial smearing effect caused by photometric redshift errors and also introducing a simple treatment for the nonlinear effects, we show that currently favored low matter density ``concordance'' $\\Lambda$CDM cosmology provides a very good fit to the estimated power. Thanks to the large volume ($\\sim 0.4 \\,h^{-3}\\,\\rm{Gpc}^3$), high clustering amplitude (linear effective bias parameter $b_{\\rm{eff}}\\sim 3\\times(0.85/\\sigma_8)$), and sufficiently high sampling density ($\\sim 3 \\times 10^{-5} \\,h^{3}\\,\\rm{Mpc}^{-3}$) the recovered power spectrum has high enough signal to noise to allow us to find weak evidence ($\\sim 2 \\sigma$ CL) for the baryonic acoustic oscillations (BAO). In case the clusters are additionally weighted by their richness the resulting power spectrum has slightly higher large-scale amplitude and smaller damping on small scales. As a result the confidence level for the BAO detection is somewhat increased: $\\sim 2.5\\sigma$. The ability to detect BAO with relatively small number of clusters is encouraging in the light of several proposed large cluster surveys. ",
        "introduction": "Since the flight of the {\\sc Cobe} satellite in the early 1990's the field of observational cosmology has witnessed extremely rapid progress which has culminated in the establishment of the Standard Model for cosmology -- the ``concordance'' model \\citep{1999Sci...284.1481B,2003ApJS..148..175S}. This progress has been largely driven by the precise measurements of the angular temperature fluctuations of the Cosmic Microwave Background (CMB) \\citep{2000ApJ...545L...5H,2002ApJ...571..604N,2003ApJS..148....1B,2003MNRAS.341L..23G,2003ApJ...591..556P}. However, in order to break several degeneracies inherent in the CMB measurements one has to complement this data with other sources of information, such as the measurements of the SNe~Ia luminosity distances \\citep{1998AJ....116.1009R,1999ApJ...517..565P} or with the measurements of the large-scale structure (LSS) of the Universe as traced by galaxies or galaxy clusters.  The simplest descriptor one can extract from the LSS measurements is the matter power spectrum. The broad-band shape of this spectrum is sensitive to the shape parameter $\\Gamma = \\Omega_mh$, and thus is useful in helping to establish the currently favored low matter density ``concordance'' model, as well as the amount of baryons in relation to the total matter $f_b=\\Omega_b/\\Omega_m$. Currently the two largest galaxy redshift surveys are the 2dF Galaxy Redshift Survey\\footnote{http://www.mso.anu.edu.au/2dFGRS/} (2dFGRS) and the Sloan Digital Sky Survey\\footnote{http://www.sdss.org/} (SDSS) with its latest data release 5 (DR5), providing redshifts to $\\sim 220,000$ and $\\sim 675,000$ galaxies, respectively. The SDSS galaxy sample consists of two broad classes: (i) the MAIN galaxy sample, reaching redshifts of $z \\sim 0.25$; (ii) the Luminous Red Galaxy (LRG) sample covering redshift range $z \\sim 0.15 - 0.5$. Some other characteristics of these surveys are listed in Table \\ref{tab}.  \\begin{table*} \\caption{Some characteristics of the 2dFGRS, SDSS DR5, and maxBCG samples.} \\label{tab} \\begin{tabular}{l|l|l|l|l|l} Survey & Number of & Sky area & Redshift & Comoving volume & Effective bias parameter wrt to\\\\ & objects & (deg$^2$) & coverage & ($h^{-3}\\,\\rm{Gpc}^3$) & the model with $\\sigma_8=0.85$\\\\ \\hline \\hline 2dFGRS & $\\sim 220,000$ & $\\sim 1200$ & $\\lesssim 0.3$ & $\\sim 0.07$ & $\\sim 1.0$ \\\\ \\hline SDSS DR5 & $\\sim 675,000$ & $\\sim 5700$ & MAIN: $\\lesssim 0.25$ & MAIN: $\\sim 0.2$ & MAIN: $\\sim 1.2$\\\\ & & & LRG: $\\sim 0.15 - 0.5$ & LRG: $\\sim 1.3$ & LRG: $\\sim 2.0$\\\\ \\hline maxBCG & $\\sim 13,800$ & $\\sim 7000$ & $ 0.1 - 0.3$ & $\\sim 0.4$ & $\\sim 3$ \\end{tabular} \\end{table*}  In addition to the precise measurement of the broad-band shape of the power spectrum, SDSS LRG and 2dFGRS samples have proven useful in allowing the detection of theoretically predicted oscillatory features in the spectrum. The source of these spectral fluctuations is the same as that giving rise to the prominent peak structure in the CMB angular power spectrum -- acoustic oscillations in the tightly coupled baryon-photon fluid prior to the epoch of recombination \\citep{1970Ap&SS...7....3S,1970ApJ...162..815P}. However, the relative level of fluctuations in the matter power spectrum is strongly reduced compared to the corresponding fluctuations in the CMB angular spectrum, owing to the fact that most of the matter in the Universe is provided by the cold dark matter (CDM) that does not participate in acoustic oscillations. After recombination, or more precisely after the end of the so-called ``drag-epoch'' which happens at redshifts of a few hundred, as the baryons are released from the supporting pressure of the photon gas they start to fall back to the CDM potential wells that have started to grow already since the matter-radiation equality. Since the amount of baryons in the total matter budget is not completely negligible -- baryons make up roughly one fifth -- the acoustic structure imprinted onto the baryonic component also gets transferred to the spatial distribution of the CDM. After baryonic and CDM components have relaxed one ends up with $\\sim 5\\%$ relative fluctuations in the total mater power spectrum.  In order to observe these relatively small fluctuations one needs to have: \\begin{enumerate} \\item Large survey volumes, to reduce cosmic variance; \\item High enough spatial sampling density, to decrease discreteness noise. \\end{enumerate} Evidently, for the fixed telescope time there is a trade-off between these two requirements. Amongst the available spectroscopic samples the optimum is currently best met by the SDSS LRGs, and indeed, at the beginning of 2005 the detection of the acoustic peak in the spatial two-point correlation function was announced by \\citet{2005ApJ...633..560E}. At the same time the final power spectrum measurements of the 2dFGRS also revealed the existence of the acoustic features \\citep{2005MNRAS.362..505C}. Baryonic acoustic oscillations (BAO) in the power spectrum of the SDSS LRG (DR4) sample were found by \\citet{2006A&A...449..891H}\\footnote{For an independent comparison see \\citet{2008PThPh.120..609Y}.} and the corresponding cosmological parameter estimation was carried out in \\citet{2006A&A...459..375H}. Also, several more recent detections of BAO have been reported: \\citet{2006astro.ph..5302P} and \\citet{2007MNRAS.374.1527B} used SDSS photometric LRG sample, \\citet{2006PhRvD..74l3507T,2008ApJ...676..889O,2009MNRAS.393.1183C,2008arXiv0807.3551G,2009ApJ...696L..93M,2009arXiv0901.2570S,2009arXiv0908.2598K} analyzed spectroscopic LRG sample, \\citet{2007ApJ...657..645P} carried out a combined analysis of the SDSS MAIN and LRG spectroscopic galaxy samples, and finally, \\citet{2007MNRAS.381.1053P,2009arXiv0907.1660P} analyzed SDSS LRG, MAIN and 2dFGRS samples. In the recent paper by \\citet{2009PhRvD..79f3512N} a weak detection of the BAO damping is reported.  The usefulness of BAO arises from the fact that it can provide us with a standard ruler in the form of the size of the sound horizon at decoupling\\footnote{To be more precise again, at the end of the ``drag-epoch''. In the ``concordance'' $\\Lambda$CDM model the sound horizon at the ``drag-epoch'' is $\\sim 5\\%$ larger compared to the one at recombination.}, enabling us to carry out a purely geometric cosmological test: comparing the apparent size of the ruler along and perpendicular to the line of sight with the physical size of the sound horizon (which can be well calibrated from the CMB data) one is able to find Hubble parameter $H(z)$ and angular diameter distance $d_A(z)$ corresponding to the redshift $z$ (e.g. \\citealt{2003PhRvD..68f3004H}). Obviously, having determined $H(z)$ and $d_A(z)$ one can put constraints on the dark energy (DE) equation of state parameter $w_{DE}(z)$. In that respect radial modes, which enable us to find $H(z)$, are more useful since $H(z)$ is related to $w_{DE}(z)$ through a single integration whereas $d_A(z)$, which is given by the angular modes, involves double integration over $w_{DE}(z)$. Thus one would certainly benefit from accurate redshift information.  If one wishes to use the BAO signal as a precise standard ruler there are a few complications one has to overcome: \\begin{enumerate} \\item As the density field goes nonlinear couplings between Fourier modes (and thus the resulting transfer of power) will modify the BAO signal, leading to the damping of the oscillations; \\item In order to fully exploit the information in galaxy/cluster surveys one has to understand the relation of these objects to the underlying matter density field, i.e. one has to have a realistic model for biasing. \\end{enumerate} These two complications can be fully addressed only through costly N-body simulations. There are several papers that have investigated the detectability and possible systematics of BAO extraction using the numerical simulations e.g. \\citet{1999MNRAS.304..851M,2005Natur.435..629S,2005ApJ...633..575S,2007APh....26..351H,2007A&A...462....7K,2007astro.ph..2543A,2007astro.ph..3620S}. However, significantly more work towards these directions is probably needed before we can harness the full power provided by the potentially very clean geometric tool in the form of BAO in the matter power spectrum.  It is worth pointing out that actually one of the first possible hints for the existence of BAO in the matter power spectrum came from the analysis of the spatial clustering of the Abell/ACO clusters \\citep{2001ApJ...555...68M}.\\footnote{Interesting features in the spatial two-point correlation function of the Abell/ACO cluster sample were also revealed by \\citet{1997MNRAS.289..801E}.} The power spectrum measurement in \\citet{2001ApJ...555...68M} used $637$ clusters with richness $R \\geq 1$. In this paper we analyze the maxBCG cluster sample \\citep{2007astro.ph..1265K} which is currently by far the largest cluster catalog available, spanning a redshift range $z=0.1-0.3$, containing $13,823$ objects, and covering $\\sim 0.4 \\,h^{-3}\\,\\rm{Gpc}^3$ of comoving volume (see Table \\ref{tab} for comparison with 2dFGRS and SDSS DR5). The biggest advantage of using galaxy clusters instead of the typical, i.e. $L_*$-galaxies, is the fact that their spatial clustering signal is strongly amplified with respect to the clustering of the underlying matter: bias parameter $b$ often reaching values $\\sim 3$ and above. \\begin{figure} \\centering \\includegraphics[width=\\plotwd]{Figs/effvol.eps} \\caption{The effective volume (see Eq. (\\ref{eq2})) of the maxBCG cluster sample in comparison to the 2dFGRS and SDSS MAIN and LRG samples from data release five (DR5).} \\label{fig1} \\end{figure} In Fig. \\ref{fig1} we have compared the effective volume (see Eq. (\\ref{eq2})) of the maxBCG survey with the two other surveys that have lead to the detection of BAO: 2dFGRS and SDSS LRG. In addition we have also shown the effective volume of the SDSS MAIN galaxy sample. Under the assumption of Gaussianity the power spectrum measurement errors $\\Delta P$ are simply found as: $\\Delta P/P = 2/V_{\\rm{eff}}/V_k$, where $V_k$ is the volume of the shell in $k$-space over which the angular average is taken, i.e. $V_k = 4 \\pi k^2 \\Delta k/(2\\pi)^3$. It is no surprise that the SDSS LRG sample is currently unbeatable. However, we notice that maxBCG sample should provide better measurement of the power spectrum on scales larger than $k \\sim 0.1 \\,h\\,\\rm{Mpc}^{-1}$ compared to the 2dFGRS. The rapid fall of the effective volume of the maxBCG sample on smaller scales is caused by the photometric redshift errors of $\\delta z \\sim 0.01$ as estimated by \\citet{2007astro.ph..1265K}.  Thus one would expect the maxBCG sample to reveal acoustic features. As our further analysis shows this indeed turns out to be the case, albeit with a relatively modest confidence level.  Our paper is organized as follows. In Section 2 we give a brief description of the maxBCG catalog with the corresponding selection effects. The core part of this paper, which is devoted to the power spectrum analysis, is given in Section 3. In Section 4 we test the stability of our spectrum measurements by relaxing several assumptions made in the original analysis. Finally, Section 5 contains our conclusions. ",
        "conclusions": "We have calculated the redshift-space power spectrum of the maxBCG cluster sample \\citep{2007astro.ph..1265K} -- currently the largest cluster catalog in existence. After correcting for the radial smoothing caused by the photometric redshift errors, treating the convolution with a survey window, and introducing a simple model for the nonlinear effects we show that ``concordance'' $\\Lambda$CDM model is capable of providing very good fit to the data. Moreover, we are able to find weak evidence for the acoustic features in the spatial clustering pattern of the maxBCG sample. If radial smoothing scale $\\sigma$, effective bias parameter $b_{\\rm{eff}}$, and nonlinear distortion parameter $q$ are treated as completely free parameters the model with acoustic oscillations is favored by $2.2\\sigma$ over the corresponding ``smooth'' model without any oscillatory behavior. The evidence for BAO weakens somewhat ($2.0\\sigma$ in case of Monte Carlo errors), and we are also left with some ``extra power'' on large scales, if we fix the smoothing scale $\\sigma$ to $30\\,h^{-1}\\,\\rm{Mpc}$ as suggested by the photometric redshift errors $\\delta z = 0.01$ estimated in \\citet{2007astro.ph..1265K}. We note that these photometric errors actually apply only for the very brightest cluster members and in reality the scatter for the whole sample might be larger. In addition the cluster finding algorithm itself might introduce extra smoothing/overmerging along the line of sight. Due to these uncertainties we have decided not to carry out any cosmological parameter study in this paper, as any untreated systematics will be immediately propagated to the parameter estimates, leading to biased results.  We have also carried out power spectrum measurement by applying additional weights that are proportional to the richness of clusters. This way we give more statistical weight to systems with a higher mass. As a result we find slightly higher large-scale clustering amplitude and mildly weaker small-scale damping of the spectrum. The latter effect is probably caused by tighter photometric redshift errors achievable for richer clusters. Larger clustering amplitude along with smaller damping leads to tighter errorbars on power spectrum measurement. Compared to the initial weighting the confidence level for the BAO detection is increased from $2.2 \\sigma$ to $2.5 \\sigma$.\\footnote{This corresponds to the case where the spatial smoothing scale due to photometric redshift errors was treated as a free parameter.}  In our paper we have estimated power spectrum errors using three different methods: (i) FKP errors, (ii) ``jackknife'' errors, obtained by dividing the survey into $5\\times 5 \\times 3$ chunks containing equal number of galaxies, (iii) Monte Carlo errors found from 1000 mock catalogs. All of these errors turn out to agree very well with each other. We have tested the stability of our power spectrum measurements by relaxing various assumptions made during the main part of our analysis. The results of this study, which were presented in Section \\ref{sec4}, demonstrate the reliability of the original analysis.  The detectability of BAO with $\\sim 10^4$ galaxy clusters is very encouraging result keeping in mind the future cluster surveys, such as the ones based on the measurement of the thermal Sunyaev-Zeldovich (SZ) effect \\citep{1972CoASP...4..173S}, e.g. {\\sc Spt}\\footnote{http://spt.uchicago.edu/}, {\\sc Planck}\\footnote{www.rssd.esa.int/Planck/}, or the proposed $100,000$-cluster X-ray survey eROSITA\\footnote{www.mpe.mpg.de/erosita/MDD-6.pdf}. On the upper panel of Fig. \\ref{newfig4} we have compared the effective volume of the maxBCG cluster sample with the SPT-like SZ survey covering one octant of the sky with the total of $\\sim 20,000$ galaxy clusters, and with the eROSITA-type X-ray survey over $50\\%$ of the sky yielding $\\sim 100,000$ clusters. We have plotted two distinct scenarios: (i) the optimistic case where one has the spectroscopic redshifts for all the clusters, (ii) the case with the photo-z errors comparable to the accuracy obtained for the maxBCG clusters, i.e. $\\delta z \\simeq 0.01$. Thus, the eROSITA-like cluster survey with spectroscopic cluster redshifts has approximately $50$ times larger effective volume compared to the maxBCG sample at $k \\sim 0.1 \\,h\\,\\rm{Mpc}^{-1}$ implying $\\sim 7$ times smaller errors for $P(k)$. If only photometric redshifts with  $\\delta z \\simeq 0.01$ are available the effective volume increases by a factor of $\\sim 12$ (which is also the case for the SPT-like survey with spectroscopic redshifts) leading to $\\sim 3.5$ times decrease in power spectrum errors. The lower panel of Fig. \\ref{newfig4} shows the contours on the wavenumber--redshift plane where $nP=3$, which is a typical target value of signal to noise for the future BAO surveys involving galaxies (see e.g. \\citealt{2006astro.ph..9591A}.) Below those lines the sampling density is high enough for the shot noise not to contribute significantly. The curve for the maxBCG sample is restricted to a redshift range $z=0.1-0.3$. It is evident that except for the lowest redshifts, i.e. $z \\lesssim 0.3$, those future cluster surveys will have significant levels of discreteness noise, and thus in terms of the raw performance of measuring the BAO only can not compete with a dedicated galaxy surveys.  In reality the main target of these cluster surveys is to map the evolution of the cluster number density as a function of redshift, since this quantity is very sensitive to the amplitude and to the growth rate of the density perturbations, and as such it provides a powerful tool to constrain the properties of DE. The possibility to also find the clustering power spectrum can be seen as a useful byproduct of these surveys that comes essentially ``for free''. However, both cluster number count study and spatial clustering analysis needs as an input estimates for the cluster redshifts, with the former probably doing relatively well with rather poor redshift estimates (e.g. photometric redshifts). This is a rather complicated task since one arguably needs of order $10$ galaxy redshifts per cluster to reliably infer cluster redshift. In that respect for the measurement of BAO it is certainly less costly to sample the cosmic density field using e.g. LRGs or blue emission line galaxies. The latter type of objects are targets for the currently ongoing WiggleZ project \\citep{2007astro.ph..1876G} at AAO that attempts to measure redshifts for $400,000$ objects over $\\sim 1000$ deg$^2$ and cover the redshift range $0.5< z < 1$. There is also a {\\sc Wfmos} instrument (see e.g. \\citealt{2005A&G....46e..26B}) construction planned for the Gemini and Subaru observatories that will hopefully be completed in 2012. This new multi-object spectrograph will be able to measure the spectra of $\\sim 5000$ objects at the same time. There are plans to perform a wide field ($\\sim 2000$ deg$^2$) redshift survey giving spectra for $\\sim 2 \\times 10^6$ galaxies up to redshifts of $z \\sim 1.3$ together with a narrower ($\\sim 200$ deg$^2$) and deeper ($z \\sim 2 - 3$) survey with a yield of $\\sim 5 \\times 10^5$ galaxies. For the more distant future (around 2020) one would expect a superb measurement of the matter power spectrum using $\\sim 10^9$ galaxy redshifts obtainable via the 21 cm measurements by the Square Kilometre Array ({\\sc Ska})\\footnote{http://www.skatelescope.org/} (e.g. \\citealt{2004NewAR..48.1063B,2005MNRAS.360...27A}). However, before {\\sc Wfmos} and {\\sc Ska} become available several wide area imaging surveys, such as Pan-STARRS\\footnote{pan-starrs.ifa.hawaii.edu/} or Dark Energy Survey\\footnote{https://www.darkenergysurvey.org/} ({\\sc Des}), will be performed. Although these photometric surveys lack accurate redshift information, the huge number of objects detectable over large sky areas significantly compensate this shortcoming, allowing us to obtain a highly competitive measurement of BAO \\citep{2005MNRAS.363.1329B}.  NOTE: \\\\ Several months after the first version of this work was submitted a paper appeared \\citep{2009ApJ...692..265E} where the authors perform correlation function analysis using the same maxBCG cluster sample. They also find weak evidence for the BAO, however their confidence level for the BAO detection ($1.4- 1.7\\, \\sigma$) is somewhat lower than we typically find in our study ($1.7-2.2 \\sigma$). It is important to point out that similar to our work \\citet{2009ApJ...692..265E} consider only isotropised quantities. Due to photo-z errors the real clustering pattern is significantly anisotropic and it is clear that isotropization looses quite some information. Here one should realize that in general the information loss for the case of the power spectrum somewhat differs from the information loss in the case of the correlation function, or equivalently, isotropised power spectrum and correlation function do not form a Fourier transform pair (only the full anisotropic quantities are the Fourier transforms of each other). Taking this into account one should not \\emph{a priori} expect similar results for the BAO detection.  It was somewhat surprising to find out that \\citet{2009ApJ...692..265E} use only five times more random points to build their unclustered catalogs. As the cluster sample itself is very sparse this surely leads to a very noisy reference level with respect to what the fluctuations are measured, leaving serious doubts on the convergence of their results.  \\citet{2009ApJ...692..265E} estimate the correlation function covariance matrix in two different ways: (i) ``jackknife'' method, (ii) simple transformation of the power spectrum covariance. In their ``jackknife'' method the subsamples are generated by each time randomly removing $1/1000$th of the clusters. This is a valid method only as long as the clusters are independent. However, due to the large-scale structure the clusters are not independent, and thus one should use methods which are able to cope better with the situation in hand. One method of dealing with dependent data is to replace the usual bootstrap/jackknife method with the block bootstrap/jackknife scheme (see e.g. \\citealt{lahiri.book}). This is indeed the reason why in large-scale structure studies the ``jackknife'' method is usually implemented by removing whole contiguous spatial chunks of data, instead of just some random set of points. By removing whole spatial regions one keeps the correlations inside the regions. If one removes random selection of points those correlations are lost, which leads to an underestimation of the off-diagonal elements in the power spectrum and correlation function covariance matrices. This might explain why their ``jackknife'' error model gives weaker confidence level for the BAO detection compared to our case. The other method of estimating covariance matrix is actually equivalent to our diagonal FKP covariance matrix case. Indeed, their Eq. (17) assumes that power spectrum covariance matrix is diagonal. \\footnote {In case of the non-diagonal power spectrum covariance matrix this equation should have a double integral over the wavenumbers, instead.} Looking at their Table 4 the relevant entry $\\Delta \\chi^2=2.4$ is quite similar to what we find, $\\Delta \\chi^2=2.8$ (see Table \\ref{tab2}), and thus is in a reasonable agreement with our FKP covariance matrix case."
    },
    "0910/0910.4458_arXiv.txt": {
        "abstract": "We report extensive spectroscopic and differential photometric $BV\\!RI$ observations of the active, detached, 1.309-day double-lined eclipsing binary IM~Vir, composed of a G7-type primary and a K7 secondary. With these observations we derive accurate absolute masses and radii of $M_1 = 0.981 \\pm 0.012~M_{\\sun}$, $M_2 = 0.6644 \\pm 0.0048~M_{\\sun}$, $R_1 = 1.061 \\pm 0.016~R_{\\sun}$, and $R_2 = 0.681 \\pm 0.013~R_{\\sun}$ for the primary and secondary, with relative errors under 2\\%. The effective temperatures are $5570 \\pm 100$~K and $4250 \\pm 130$~K. The significant difference in mass makes this a favorable case for comparison with stellar evolution theory. We find that both stars are larger than the models predict, by 3.7\\% for the primary and 7.5\\% for the secondary, as well as cooler than expected, by 100~K and 150~K, respectively. These discrepancies are in line with previously reported differences in low-mass stars, and are believed to be caused by chromospheric activity, which is not accounted for in current models. The effect is not confined to low-mass stars: the rapidly-rotating primary of IM~Vir joins the growing list of objects of near-solar mass (but still with convective envelopes) that show similar anomalies. The comparison with the models suggests an age of 2.4~Gyr for the system, and a metallicity of [Fe/H] $\\approx -0.3$ that is consistent with other indications, but requires confirmation. ",
        "introduction": "\\label{sec:introduction}  Our knowledge of stellar structure and evolution rests heavily on the comparison between theory and observation. Double-lined eclipsing binaries (hereafter EBs) have long been at the center of this process, since they allow the mass -- the most fundamental of all stellar properties -- as well as the radius to be determined to very high precision (and accuracy), often as good as 1\\%, independently of the distance and independently of any calibrations. Such high precision enables stringent tests of theory, as described, e.g., by \\cite{Andersen1991}.  In the last decade or so it has become clear that stars in the lower main-sequence show significant discrepancies when compared to standard models. Studies of several key systems have shown unambiguously that the radii predicted by the models are systematically up to $\\sim$10\\% too small, and the temperatures $\\sim$5\\% too high \\citep[see][and references therein]{Ribas2008}. The deviations are commonly attributed to the high level of chromospheric activity present in these systems. Orbital periods are typically short, and as a result tidal forces tend to synchronize the components' rotation with the orbital motion. The rapid rotation, in turn, induces the activity, which can manifest itself in the form of copious X-ray emission, flaring, H$\\alpha$ and Ca~II H and K emission, spottedness, and other effects.  According to the recent compilation by \\cite{Torresetal:09}, among the eclipsing binaries with at least one component below 0.8~M$_{\\odot}$ only five have reliable mass and radius determinations with relative errors that are smaller than 3\\%. These are, in order of decreasing mass, UV~Psc \\citep{Popper:97}, GU~Boo \\citep{LopezMorales2005}, YY~Gem \\citep{Torres2002}, CU~Cnc \\citep{Ribas2003}, and CM~Dra \\citep{Morales2009}. A number of other similar systems are known, but are not yet at the same level of precision \\emph{and} accuracy. For further progress it is therefore important to either improve the determinations in the latter cases, or to find new ones where those conditions are met. The eclipsing system reported here is one of these, which in addition presents the largest difference in mass between the components, providing in principle extra leverage for the comparison with models.  IM~Virginis (also HD~111487, 1E~1247.0$-$0548, $\\alpha = 12^{\\rm h}49^{\\rm m}38.\\!^{\\rm s}70$, $\\delta = -6^{\\circ}04'44.\\!''9$, J2000.0; \\ion{G7}{5}, $V = 9.57$) was detected as an X-ray source with the \\textit{Einstein} Observatory by \\cite{Helfand1982}. The radial-velocity variability was found by \\cite{Silva1987}, and subsequent spectroscopic and photometric studies carried out by \\cite{Marschall1988, Marschall1989} confirmed IM~Vir to be both double-lined and eclipsing, and to be composed of a G7 star and a late-K or early-M star, thus piquing our interest. The orbital period was estimated as 1.3085 days.  Very little data on this binary have been published since, other than sparse photometry and occasional reports on the chromospheric activity and X-ray flaring \\citep{Strassmeier1993, Pandey2008}. We present here extensive spectroscopic and differential photometric measurements that allow us to determine its fundamental properties with very high accuracy. The observations are presented in \\S\\thinspace\\ref{sec:spectroscopy} and \\S\\thinspace\\ref{sec:lc_obs}, followed by a detailed light-curve analysis in \\S\\thinspace\\ref{sec:lc_analysis} accounting for the presence of spots. The absolute dimensions are derived in \\S\\thinspace\\ref{sec:absolute} with careful consideration of potential sources of systematic error, essential for the results to be useful. The comparison with models of stellar evolution is discussed in \\S\\thinspace\\ref{sec:discussion}, and we summarize our conclusions in \\S\\thinspace\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  Our extensive spectroscopic and photometric observations of IM~Vir have enabled us to determine highly accurate values for the absolute masses and radii of both stars to better than 2\\%, as well as accurate temperatures (Table~\\ref{tab:absolute}). This eclipsing system now joins the ranks of those with the best determined properties \\citep[see][]{Torresetal:09}. The primary is of solar type, and the secondary is a late K dwarf. The very different masses provide increased discriminating power for testing models of stellar evolution. We find evidence that both components, which have convective envelopes and are rapidly rotating, show discrepancies in their radii and temperatures compared to calculations by \\cite{Baraffe1998}, similar to those reported for other low-mass stars. The predicted radii are too small by 3.7\\% and 7.5\\% for the primary and secondary, and the theoretical temperatures are too high by 100~K and 150~K, respectively. These effects are ascribed to chromospheric activity, for which there is abundant evidence in this system in the form of X-ray, \\ion{Ca}{2}~H and K, and H$\\alpha$ emission, X-ray flaring, and spottedness seen directly in the light curves. The larger effect observed for the secondary is consistent with the larger fraction of that star being covered by spots (a factor of two difference). The fact that the near-solar mass primary is also affected supports recent findings for other binaries indicating that the impact of activity on the structure of stars is not limited to the M dwarfs, as previously thought, but most likely extends to all stars with convective envelopes, reaching up to masses near or larger than that of the Sun.  Current stellar evolution models do not account for these effects. Some progress in this area has been made to incorporate magnetic activity in the theoretical calculations \\citep[e.g.,][]{Dantona:00, Mullan:01, Chabrier:07}, and initial results are very encouraging.  A crucial ingredient for the comparison with models is the chemical composition. The best-fit models for IM~Vir indicate an age of $2.4 \\pm 0.5$~Gyr and a metallicity of [Fe/H] $= -0.28 \\pm 0.10$. The uncertainties here represent the sensitivity to systematic errors of $\\pm$100~K in the effective temperatures, but exclude unquantified systematics in the \\cite{Baraffe1998} models themselves, which are difficult to evaluate. While the metallicity prediction seems to be in good agreement with our rough spectroscopic abundance estimate in \\S\\thinspace\\ref{sec:spectroscopy} and that of \\cite{Dall:07}, as well as photometric estimates for both stars, a proper detailed spectroscopic analysis of both components is still lacking, and is essential to validate the comparison.  Such determinations are notoriously difficult in double-lined binaries, which is why relatively few have them. In IM~Vir this task would be challenging for two reasons: the secondary is very faint compared to the primary, and it is of late spectral type. Abundance determinations for late-type stars are still problematic due to shortcomings in the model atmospheres. However, IM~Vir offers a unique opportunity because the secondary eclipse is total. High S/N ratio spectra taken with a sufficiently large telescope during the 28-minute totality phase would be of the primary only, and can be analyzed with standard techniques. The uncertainty in this observing window is estimated to be approximately $\\pm$5 minutes, based on the uncertainties in the geometric light-curve parameters."
    },
    "0910/0910.4202_arXiv.txt": {
        "abstract": "We study a generalized version of Chaplygin gas as unified model of dark matter and dark energy. Using realistic theoretical models and the currently available observational data from the age of the universe, the expansion history based on the type Ia supernovae, the matter power spectrum, the cosmic microwave background radiation anisotropy power spectra, and the perturbation growth factor we put the unified model under observational test. As the model has only two free parameters in the {\\it flat} Friedmann background [$\\Lambda$CDM (cold dark matter) model has only one free parameter] we show that the model is already tightly constrained by currently available observations. The only parameter space extremely close to the $\\Lambda$CDM model is allowed in this unified model. ",
        "introduction": "A unified dark matter and dark energy model motivated by the (generalized) Chaplygin gas has been introduced in the literature \\cite{Kamenshchik-etal-2001,Bilic-etal-2002,Bento-etal-2002}, and has been extensively investigated in the cosmological studies \\cite{Bean-Dore-2003,Bento-etal-2003,Carturan-Finelli-2003,Avelino-etal-2003, Amendola-etal-2003,Bento-etal-2004,Silva-Bertolami-2003,Bertolami-etal-2004, Cunha-etal-2004,Zhu-2004,Dev-etal-2003, Gorini-etal-2003,Makler-etal-2003a, Makler-etal-2003b,Alam-etal-2003,Colistete-etal-2004,Dev-etal-2004, Multamaki-etal-2004,Sandvik-etal-2004,Beca-etal-2003,Alcaniz-Lima-2005, Colistete-Fabris-2005,Bento-etal-2005,Biesiada-etal-2005,Gong-2005, Lazkoz-etal-2005,Bertolami-Silva-2006,Zhang-Zhu-2006,Zhang-etal-2006, Wu-Yu-2007a,Wu-Yu-2007b,Guo-Zhang-2007,Barreiro-etal-2008,Gorini-etal-2008, Fabris-etal-2008,Li-etal-2009,Urakawa-Kobayashi-2009}. Many constraints on the Chaplygin gas model parameters have been placed (predicted) based on current (upcoming) astronomical observations such as type Ia supernovae (SNIa) \\cite{Bean-Dore-2003,Amendola-etal-2003,Avelino-etal-2003,Makler-etal-2003a, Makler-etal-2003b,Beca-etal-2003,Bertolami-etal-2004,Colistete-etal-2004, Dev-etal-2004,Cunha-etal-2004,Zhu-2004,Colistete-Fabris-2005,Bento-etal-2005, Biesiada-etal-2005,Gong-2005,Zhang-etal-2006,Guo-Zhang-2007,Alcaniz-Lima-2005, Lazkoz-etal-2005,Wu-Yu-2007b,Alam-etal-2003,Silva-Bertolami-2003}, large-scale structure \\cite{Bean-Dore-2003,Avelino-etal-2003,Sandvik-etal-2004,Beca-etal-2003, Multamaki-etal-2004,Zhang-etal-2006,Gorini-etal-2008,Fabris-etal-2008, Urakawa-Kobayashi-2009}, cosmic microwave background radiation (CMB) \\cite{Bean-Dore-2003,Bento-etal-2003,Carturan-Finelli-2003,Amendola-etal-2003, Zhang-etal-2006,Barreiro-etal-2008,Urakawa-Kobayashi-2009}, gravitational lensing \\cite{Dev-etal-2003,Dev-etal-2004,Silva-Bertolami-2003,Makler-etal-2003b}, gamma-ray bursts \\cite{Bertolami-Silva-2006}, X-ray luminosity of galaxy clusters \\cite{Makler-etal-2003b,Cunha-etal-2004,Zhu-2004}, look-back time-redshift data \\cite{Gong-2005,Li-etal-2009}, angular size-redshift data \\cite{Alcaniz-Lima-2005}, Hubble parameter-redshift data \\cite{Wu-Yu-2007a,Wu-Yu-2007b}, Fanaroff-Riley type IIb radio galaxies \\cite{Makler-etal-2003b,Zhu-2004}, and so on.  It is known that such a unified model is constrained by observations so that parameter space close to the conventional $\\Lambda$CDM models is allowed. But current status of the model in parameter space different from close-to-$\\Lambda$CDM has been unclear. The distance-redshift relation, the age of the universe, abundances of light elements, the large-scale matter density and velocity power spectrum, the CMB temperature and polarization anisotropy power spectra, and the perturbation growth factor can be regarded as the main pillars of modern cosmology where theories meet with observations. Here, we investigate the unified model by comparing realistic theoretical predictions with currently available observations. Our conclusions can be found in section \\ref{sec:Discussion}. ",
        "conclusions": "\\label{sec:Discussion}  We have systematically studied observational consequences of GCG unified dark matter and dark energy model. In the background Friedmann world model the basic requirement of minimum age of the universe together with the SNIa luminosity-redshift data already leaves only small parameter space close to the $\\Lambda$CDM model, see Fig.\\ \\ref{fig:SN-mcmc-GCG}. In the perturbation study, we have shown that the matter power spectrum and the CMB power spectra are mutually exclusive except for narrow region close to the $\\Lambda$CDM model, see Figs.\\ \\ref{fig:power-spectra-1}, \\ref{fig:power-spectra-2}, and \\ref{fig:power-spectra-3}. A tight constraint on $\\alpha$-parameter in that allowed region can be obtained from the (baryonic) matter power spectrum as $-5\\times 10^{-5} \\le \\alpha \\le 10^{-4}$, see Figs.\\ \\ref{fig:power-spectra-2} and \\ref{fig:power-spectra-2-X}. More realistic constraint can be found in Fig.\\ \\ref{fig:matter-PS-MCMC} and Table \\ref{tab:GCG_constraints} as $|\\alpha| \\lesssim 10^{-4}$ from the SDSS DR7 LRG power spectrum and $|\\alpha| \\lesssim 10^{-5}$ from the $\\Lambda\\textrm{CDM}$ mock power spectrum data. Although this allowed region is wider than what was concluded based on the GCG power spectrum in \\cite{Sandvik-etal-2004} (see Fig.\\ \\ref{fig:power-spectra-2-X}), the constraint based on the baryon power spectrum is still severe enough so that the observationally allowed GCG model can be regarded as extremely close to the $\\Lambda$CDM model, thus, effectively indistinguishable from the $\\Lambda$CDM model. Although named as a unified model the GCG model has two free parameters, $\\alpha$ and $w_{X0}$, whereas the $\\Lambda$CDM has only one free parameter $\\Omega_{\\Lambda 0}$. From the perspective of more wider parameter space of our GCG model, it is remarkable to notice the distinguished success and its uniqueness of the $\\Lambda$CDM model.   \\subsection*"
    },
    "0910/0910.2607.txt": {
        "abstract": "Radio-bright BL Lacertae objects (BLOs) are typically very variable and exhibit prominent flaring. We use a sample of 24 BLOs, regularly monitored at Mets\\\"ahovi Radio Observatory, to get a clear idea of their flaring behavior in the radio domain and to find possible commonalities in their variability patterns. Our goal was to compare the results given by computational time scales and the observed variability parameters determined directly from the flux curves. Also, we wanted to find out if the BLO flares adhere to the generalized shock model, which gives a schematic explanation to the physical process giving rise to the variability. We use long-term monitoring data from 4.8, 8, 14.5, 22, 37, 90 and 230 GHz, obtained mainly from University of Michigan and Mets\\\"ahovi Radio Observatories. The structure function, discrete correlation function and Lomb-Scargle periodogram time scales, calculated in a previous study, are analyzed in more detail. Also, we determine flare durations, rise and decay times, absolute and relative peak fluxes from the monitoring data. We find that radio-bright BLOs demonstrate a wide range of variability behavior, and few common denominators can be found. BLOs include sources with fast and strong variability, such as OJ 287, PKS 1749+096 and BL Lac, but also sources with more rolling fluctuations like PKS 0735+178. The most extreme flares can last for up to 13 years or have peak fluxes of approximately 12 Jy in the observer's frame. When the Doppler boosting effect is taken into account, the peak flux of a flare does not depend on the duration of the flare. A rough analysis of the time lags and peak flux evolution indicates that, typically, BLO flares in the mm -- cm wavelengths are high-peaking, i.e., are in the adiabatic stage. Thus, the results concur with the generalized shock model, which assigns shocks travelling in the jet as the main cause for AGN variability. Comparing the computational time scales and the parameters obtained from the flux curve analysis (i.e., rise and decay times and intervals of the flares) reveals that they do have a significant correlation, albeit with large scatter. ",
        "introduction": "BL Lacertae objects (BLOs) are a relatively rare subclass of active galactic nuclei (AGN). Traditionally the defining properties of BLOs include a featureless optical spectrum, a flat radio spectrum and vigorous variability at all frequency bands \\citep[][and references therein]{stein76, kollgaard94, jannuzi94, urry95}. Most BLOs are thought to be highly beamed objects \\citep{blandford79}, which means that the relativistic jets emanating from the core are pointing very closely at our direction. This is partly the cause of the featureless optical spectrum; the non-thermal continuum emission from the jet can swamp the thermal emission, including the emission lines, from the host galaxy. However, in the case of less beamed objects, the lineless spectrum must be created by other mechanisms.  The first BLO samples were discovered in either radio \\citep{stickel91, stickel93} or X-ray surveys \\citep{gioia90, stocke91}. Lately, surveys at different wavelengths \\citep{londish02} and the cross-correlation of existing radio and other waveband catalogs \\citep{perlman98,caccianiga99, landt01, giommi05, turriziani07, plotkin08} have produced a large number of new members to the class. Typically the selection criteria are different from those of the pioneering surveys, which has widened the definition of a BL Lac object; hence, the variability characteristics of many of the newest BLOs are unknown due to lack of data.  In a recent paper \\citep{nieppola07}, 37 GHz fractional variability indices from Mets\\\"ahovi Radio Observatory data were calculated for 90 BLOs. All sources, for which even a crude estimation of variability could be calculated, exhibited an increase of 10 \\% of the minimum flux level at some point during the 3.5 year observation period. Almost half of them doubled their minimum flux density. However, this sample of 90 sources was only one fourth of the full Mets\\\"ahovi BLO sample; the rest are too faint to allow any variability analysis.  The variability of BLOs, like all flaring AGN, is thought to be caused by shocks forming and travelling in the jet. The origin and early development of these shocks is not yet very well understood. Once propagating downstream in the jet, their evolution is easier to model with Compton, synchrotron and adiabatic losses \\citep{mar:gear85, hughes89}. \\citet{valtaoja92_moniIII} constructed a generalized view of the Marscher \\& Gear shock model, containing a general scenario of the AGN flare behavior, to be used in comparison with observations. The generalized model describes how the shape of the shock spectrum remains unchanged as its peak moves from higher to lower frequencies. The evolutionary track of the shock consists of growth ($S_m\\propto\\nu^a_m$), plateau ($S_m\\approx\\textrm{constant}$), and decay ($S_m\\propto\\nu^b_m$) stages, where $S_m$ and $\\nu_m$ are the turnover flux and frequency for the shock spectrum, and $a$ and $b$ are model-dependent parameters. The flares can ideally be divided into two groups: i) low-peaking flares, which will reach their maximum intensity at lower frequencies than the observing frequency, and ii) high-peaking flares, which have peaked at high frequencies relative to the observing frequency and are already decaying. The observing frequency, however, is not a constant quantity, but can be chosen freely. Thus the low- and high-peaking classes are not fixed either: the same flare can be low-peaking in one frequency band and high-peaking on another. Consequently, the classes are not separated, but rather opposite ends of a continuum of cases. Low- and high-peaking flares should be distinguishable from observations. For high-peaking flares the peak fluxes and time lags are strongly frequency-dependent, the highest observing frequency peaking first with the highest peak flux. For low-peaking flares the peaks are nearly simultaneous in all frequencies, and the peak flux of the flare is not significantly dependent on the frequency.  In this work we will study the long-term radio variability of BL Lacertae objects at several frequencies. We will focus on a sample of 24 BLOs. The current number of definite and probable BLOs is over 1000 \\citep{veron06}. This means that our sample is not a representative cross-section for the whole population, but rather represents only the rare, radio luminous BLOs. The bases of our work are the extensive databases of Mets\\\"ahovi and University of Michigan Radio Observatories at frequencies 4.8, 8, 14.5 GHz (UMRAO) and 22 and 37 GHz (Mets\\\"ahovi).  We will study the variability of our BLO sample from two points of view: computational time scales obtained using statistical analysis methods and observed parameters of the flares, e.g. the duration, rise and decay times, and absolute and relative peak fluxes. The goal is to determine how well the computational time scales correspond to the behavior we observe in the source, as well as to gain a deeper understanding of what kind of variability can be expected from radio-luminous BLOs, when monitored for tens of years. We are also interested in how well the BLO flares adhere to the generalized shock model of \\citet{valtaoja92_moniIII}, and whether the flares are mostly high- or low-peaking. A similar study is performed on a larger AGN sample, including radio-loud quasars, in an accompanying paper \\citep{hovatta08b}.  In \\S\\ref{sam} we will present our sample and data. We briefly describe the methods we used in \\S\\ref{met}. In \\S\\S\\ref{ts} and \\ref{flares} we discuss the time scales and the observed flux curves, respectively. In \\S\\ref{corr} we examine the correspondence between the time scales and observed flaring, and describe the behavior of the sample sources individually. We will finish with a discussion and conclusions in \\S\\S\\ref{dis} and \\ref{con}, respectively. Throughout this paper, we assume $H_{0}=72\\,\\textrm{km}\\,\\textrm{s}^{-1}\\,\\textrm{Mpc}^{-1}$, $\\Omega_m=0.27$, and $\\Omega_{\\Lambda}=0.73$. ",
        "conclusions": "\\label{con} We have studied the long-term radio variability and the flare morphology of radio-bright BL Lacertae objects. The main conclusions are as follows:  \\begin{enumerate} \\item Radio-bright BLOs exhibit a range of flaring behavior, with few common features. Especially the quasar-like objects PKS 0735+178, 1308+326 and 3C 446 have distinctively long outbursts with modest short-term variability, contrasting with the more erratic behavior of other sample sources. The long flare of 1308+326 clearly stands out from the rest of the sample even after correcting for the relativistic boosting effects. Our findings confirm the quasar-like nature of 1308+326 and 3C446 and indicate that PKS 0735+178 also has radio behavior different from typical BLOs.  \\item The median duration of a flare in a radio-bright BLO is of the order of 2.5 years, and the peak flux density typically reaches about 5 Jy at 37 GHz. On average, the decay time of the flare is 1.6 times longer than the rise time. When the Doppler boosting effect is taken into account, the peak flux of the flare does not depend on the duration of the flare, indicating that the energy release in a flare does not depend on its duration.  \\item The 45 BLO flares in our analysis confirm the generalized shock model of \\citet{valtaoja92_moniIII} with no clear, undisputed exceptions. However, very frequent sampling on several radio frequencies is needed for the accurate, observational determination of the flare components and time lags. Based on the evolution of the relative peak flux and the time lags from one frequency band to the next, we find that the BLO flares are mostly high-peaking. Probably they reach their maximum development in the mm- to sub-mm-wavelengths.  \\item The computational time scales, $T_{SF}$, $T_{DCF}$ and $T_{P}$, have a statistically significant correlation with the temporal flare parameters obtained directly from the flux curves. However, scatter is considerable, and the average deviation from one-to-one correspondence is of the order of 1 - 3 years, depending on the parameter and time scale in question.  \\end{enumerate}"
    },
    "0910/0910.5803_arXiv.txt": {
        "abstract": "We report the result of a search for the infrared counterparts of 37 planetary nebulae (PNs) and PN candidates % in the  {\\it Spitzer} Galactic  Legacy Infrared Mid-Plane Survey Extraordinaire II (GLIMPSE~II) survey.  The photometry and images of these PNs at 3.6, 4.5, 5.8, 8.0, and 24\\,$\\mu$m, taken through the Infrared Array Camera (IRAC) and the Multiband Imaging Photometer for Spitzer (MIPS), are presented. Most of these nebulae are  very red and compact in the IRAC bands, and are found to be bright and extended in the 24\\,$\\mu$m band. The infrared  morphology of these objects are compared with H$\\alpha$ images of the Macquarie-AAO-Strasbourg (MASH) and MASH~II PNs. The implications for morphological difference in different wavelengths are discussed. The IRAC data allow us to differentiate between PNs and \\ion{H}{2} regions and be able to reject non-PNs from the optical catalogue (e.g., PNG\\,352.1$-00.0$). Spectral energy distributions (SEDs) are constructed by combing the IRAC and MIPS data with existing near-, mid-, and far-IR photometry measurements.  The anomalous colors of some objects allow us to infer the presence of aromatic emission bands. These multi-wavelength data provide useful insights into the nature of different nebular components contributing to the infrared emission of PNs. % ",
        "introduction": "The traditional view of planetary nebulae (PNs) is that it is a shell of gas ionized by a hot central star and exhibits an emission-line spectrum.  The theoretical understanding that PNs descend from asymptotic giant branch (AGB) stars and should contain remnants of the dust envelope led to the realization that PNs should also be bright infrared objects \\citep{kwok82}.  This was confirmed after the {\\it Infrared Astronomical Satellite} ({\\it IRAS}) all-sky survey found that over 1000 PNs show far infrared emission \\citep{pot84, zk91}.  The typical color temperature of the dust component in PNs is between 100 and 200 K, warmer than the dust component in H~{\\sc ii} regions but cooler than the circumstellar dust envelopes of AGB stars.  This is consistent with the fact that the dust envelopes in PNs represent the dispersing remnants of the AGB envelopes \\citep{kwok90}.  Infrared emission is now recognized as a defining observational properties of PNs.  However, the location of the infrared emitting dust component remains poorly known.  Only recently has the angular resolution of mid-infrared cameras  on large ground-based telescopes with adaptive optics become high enough to image the  distribution of dust in PNs.  Space observations therefore offer a better opportunity to survey a large number of PNs.  The  Infrared Array Camera (IRAC)  on the  {\\it Spitzer Space Telescope} has an angular resolution $<2$ arcsec \\citep{fazio04} and its capability in imaging the dust component of PNs has been  clearly demonstrated by \\citet{hora04}.  Since the peak of the dust emission from PNs occurs at wavelengths $>20$ $\\mu$m, the Multiband Imaging Photometer for Spitzer \\citep[MIPS;][]{rieke04}, although offering less angular resolution than IRAC, would provide useful information on the distribution of the main cold dust component.  Based  on the MIPS observations at 24 and 70\\,$\\mu$m, \\citet{su07} found possible evidence for the presence of a debris disk around the central star of the Helix Nebula.  Other examples of PNs studied by {\\it Spitzer} includes NGC\\,2346 \\citep{su04}, NGC\\,650 \\citep{ueta06}, and a sample of PNs in the Large Magellanic Cloud \\citep[LMC;][]{hora08}.  One of the advantages of IR observations is that they are hardly affected by dust reddening and are therefore ideal to study PNs heavily obscured by interstellar dust.  From the IRAC observations, \\citet{cohen05a} identified an extremely red PN G313.3+00.3 which is optically invisible.  This raises the possibility that there are many obscured PNs in the Galaxy, for which {\\it Spitzer} IR observations could explore.  A high-spatial-resolution and high-sensitivity H$\\alpha$ survey of the southern Galactic plane has been undertaken using the Anglo-Australiam Observatory UK Schmidt Telescope (AAO/UKST) to search for very faint PNs \\citep{par99}.  This has led to the discovery of over 1000 new PNs and PN candidates (which we refer to hereafter simply as PNs).  The results of this survey was published as the Macquarie/AAO/Strasbourg H$\\alpha$ Planetary Galactic Catalogue (MASH) \\citep{parker06} and the MASH~II supplement \\citep{miszalski08}.  These new PNs are typically more obscured than known Galactic PNs and warrant further investigation at IR wavelengths.  The {\\it Spitzer} GLIMPSE (Galactic Legacy Infrared Mid-Plane Survey Extraordinaire; see Benjamin et al. 2003 and Churchwell et al. 2009) survey provides IRAC images with high resolution and sensitivity at wavelengths from 3.6 to 8.0\\,$\\mu$m, allowing us to perform high-quality infrared imaging and photometry of PNs. The first-release GLIMPSE~I data cover a Galactic region of $\\left|l\\right|=10^\\circ$--65$^\\circ$ and $\\left|b\\right|\\le1^\\circ$ (total area $\\sim220$ square degrees). Comprehensive GLIMPSE~I-based studies of PNs have been presented by some authors \\citep[e.g.][]{cohen07b,phillips08a, phillips08b,ramos08}. \\citet[][hereafter Paper~I]{kwok08} imaged 30 PNs\\footnote{Note that some objects in Paper~I are actually compact \\ion{H}{2} regions, but are referred to as `PNs' for convenience of denotation (their non-PN nature has been discussed in text).} detected in the GLIMPSE~I survey and constructed IR spectral energy distributions (SEDs) which clearly show a distinction between the photospheric, nebular bound-free, and dust emission components.  In this paper, we extend the work of Paper~I to the GLIMPSE~II survey area that covers the Galactic bulge regions. We present the photometry and images of 37 PNs detected in the GLIMPSE~II survey. These results are compared with those of the GLIMPSE~I PNs. To better explore nebular dust emission, we also take into account the 24\\,$\\mu$m data from the {\\it Spitzer} Legacy program MIPS Inner Galactic Plane Survey (MIPSGAL). This paper is organized as follows: in Section~2 we briefly introduce the GLIMPSE~II observations and the {\\it Spitzer} data used in this study; in Section~3 we describe  the data processing; in Section~4 we present the Mid-IR images and our measurement results of fluxes, sizes, and SEDs for individual PNs; the astrophysical implications of our results are discussed in Section~5; and Section~6 summarizes the main conclusions. ",
        "conclusions": "An investigation of the MASH PNs within the GLIMPSE~II field has revealed that 37 optically identified PNs have MIR counterparts. These PNs are distinguishable from the surrounding stars by their redder color and more extended angular sizes. From their infrared morphologies, we suggest that some of these objects (PNG 352.1-00.0, PNG 353.9+00.0, PNG 357.4-01.3, PNG 358.8-00.0) are unlikely to be PNs.  We present the spatially integrated flux and size measurements of these PNs at 3.6, 4.5, 5.8, 8.0, and 24\\,$\\mu$m. MIR color-color and color-magnitude diagrams are produced. Combining our results with those in other data archives, we construct the 0.82--100\\,$\\mu$m SEDs for these objects. The primary findings from analysis of these results are  summarized below.  \\begin{list}{}{}  \\item[1.]  Compared to the GLIMPSE~I PNs, the GLIMPSE~II PNs are generally more compact and have lower spatially integrated fluxes, implying that these PNs towards the Galactic bulge might be generally more distant. Nevertheless, the MIR colors of GLIMPSE~I and II PNs are very similar to each other, and are clearly different with those of stellar objects.  \\item[2.] Despite some GLIMPSE~II PNs are invisible in the 3.6--5.8\\,$\\mu$m bands and are only marginally above the IRAC 8.0\\,$\\mu$m  detection limitation, all of them are characterized with very strong 24\\,$\\mu$m emission, which is likely to originate from extended cold dust envelopes.  \\item[3.] The nebular morphologies and fluxes detected in different wavelengths allow us to separate the different components in the nebulae. The 5.8 and 8.0\\,$\\mu$m emission may be either more or less compact than their optical counterparts, respectively implying the presence of dust toruses centered on bipolar PNs or extended neutral envelopes around ionized gaseous nebulae.  \\item[4.] From the anomalous colors of some objects, we infer the contributions of AIB emissions to the IRAC fluxes.  \\item[5.] The absolute diffuse calibrations are examined by comparing the IRAC 8.0\\,$\\mu$m and MIPS 24\\,$\\mu$m to the {\\it MSX} 8.3 and 21\\,$\\mu$m fluxes. We obtain the ${\\rm IRAC}8.0/MSX8.3$ and ${\\rm MIPS}24/MSX21$ flux ratios of $0.90\\pm0.36$ and $1.19\\pm0.12$, respectively. These results are reasonably consistent with previous studies by \\citet{cohen07a}, \\citet{cohen07b}, and \\citet{cohen09} using data for \\ion{H}{2} regions and GLIMPSE~I PNs, and suggest good calibrations for these slightly resolved GLIMPSE~II PNs.  \\item[6.] We explore the MIR/radio relation for PNs. A loose correlation between the MIR (8.0\\,$\\mu$m and 24\\,$\\mu$m) and the NVSS 1.4\\,GHz fluxes is found, and can be statistically described by $q_{\\rm IR}$---the ratio between the flux densities at MIR and 1.4\\,GHz. For the GLIMPSE~II PNs, we find  $q_{8.0}=0.79\\pm0.52$ and $q_{24}=2.13\\pm0.37$.  The larger $q_{\\rm IR}$ values and the larger dispersion in the MIR/radio correlation of PNs compared to those of star-forming galaxies are attributed to their respective different underlying radiation mechanisms.   \\end{list}  Since the bulk of GLIMPSE~II PNs are located at higher Galactic latitudes ($1^\\circ\\le|b|\\le2^\\circ$), we anticipate that a large number of PNs might have been detected in the GLIMPSE~3D legacy survey, which extends the latitude coverage to $\\pm3^\\circ$. The study of GLIMPSE~3D PNs will be presented in a subsequent paper."
    },
    "0910/0910.5204_arXiv.txt": {
        "abstract": "Dans le mod\\`ele cosmologique, l'Univers est constitu\\'e principalement de mati\\`ere noire et d'\\'energie noire dont nous ne connaissons pas la nature, et pour lesquelles il n'y a pas d'explication dans le cadre du mod\\`ele standard de la physique des particules. Comment incorporer mati\\`ere et \\'energie noires dans l'ensemble des lois fondamentales\\,? Toutes les observations peuvent aussi bien \\^etre expliqu\\'ees soit par l'addition de composants de l'Univers inconnus, avec la relativit\\'e g\\'en\\'erale comme th\\'eorie de la gravitation, soit par une modification fondamentale de cette th\\'eorie. Ne serait-il pas plus simple de la modifier\\,? C'est le cas de l'hypoth\\`ese MOND (pour MOdified Newtonian Dynamics) propos\\'ee par Milgrom en 1983, qui est pleine de succ\\`es pour d\\'ecrire la cin\\'ematique et la dynamique des galaxies. Il serait cependant possible d'obtenir le m\\^eme succ\\`es gr\\^ace \\`a une nouvelle forme de mati\\`ere, la mati\\`ere noire dipolaire, tout en gardant la relativit\\'e g\\'en\\'erale comme loi de la gravitation. ",
        "introduction": "Les avanc\\'ees consid\\'erables de ces derni\\`eres ann\\'ees mettent donc en \\'evidence une grande tension dans notre vision de l'Univers\\,: celui-ci est constitu\\'e essentiellement de mati\\`ere noire et d'\\'energie noire dont nous ne connaissons pas la nature, et pour lesquelles il n'y a pas d'explication dans le cadre actuel de la physique des particules. La mati\\`ere noire, de nature exotique, pourrait \\^etre form\\'ee de particules pr\\'edites par certaines extensions du mod\\`ele standard comme dans les th\\'eories supersym\\'etriques. Mais ces particules n'ont toujours pas \\'et\\'e mises en \\'evidence de mani\\`ere directe, malgr\\'e de nombreuses exp\\'eriences ayant tent\\'e de les d\\'etecter (peut-\\^etre que le LHC va bient\\^ot nous en dire plus). Quant \\`a l'\\'energie noire, c'est un nouveau composant myst\\'erieux qui acc\\'el\\`ere l'expansion de l'Univers, et qui pourrait \\^etre la fameuse constante cosmologique $\\Lambda$ qu'Einstein avait introduite dans les \\'equations de la relativit\\'e g\\'en\\'erale, pour des raisons maintenant disparues. Le probl\\`eme est que toutes les tentatives pour interpr\\'eter $\\Lambda$ en termes de physique fondamentale ont pour l'instant \\'echou\\'e.  Mati\\`ere noire et \\'energie noire sont comme deux nuages qui se profilent \\`a l'horizon du ciel radieux de la physique. Sommes-nous \\`a l'aube d'une r\\'evolution fondamentale\\,? Plusieurs satellites, comme Euclid en Europe, ou JDEM aux \\'Etats-Unis, vont mesurer plus pr\\'ecis\\'ement les caract\\'eristiques des composantes noires, mais on peut argumenter que le probl\\`eme est plut\\^ot d'ordre th\\'eorique\\,: comment incorporer mati\\`ere et \\'energie noires dans un ensemble de lois fondamentales\\,?  Toutes les observations peuvent aussi bien \\^etre expliqu\\'ees soit par l'addition de composants inconnus, en restant dans le cadre de la relativit\\'e g\\'en\\'erale comme th\\'eorie de la gravitation, soit par une modification fondamentale de cette th\\'eorie. Ne serait-ce pas plus simple de la modifier\\,? Au XIX$^\\text{\\`eme}$ si\\`ecle l'astronome fran\\c{c}ais Le Verrier avait d\\'ecouvert une rotation anormale (dite pr\\'ecession) de l'orbite de la plan\\`ete Mercure. On pouvait l'expliquer soit par de la ``mati\\`ere noire'', en l'occurrence une nouvelle plan\\`ete int\\'erieure \\`a l'orbite de Mercure (d\\'enomm\\'ee Vulcain par Le Verrier), soit par une modification des lois de Newton. C'est cette deuxi\\`eme option qui s'est finalement av\\'er\\'ee correcte, lorsque la gravitation newtonienne a \\'et\\'e remplac\\'ee par celle d'Einstein, et que l'exc\\`es de pr\\'ecession de Mercure a \\'et\\'e expliqu\\'e par un effet relativiste. \\begin{figure}[t] \\centering{\\includegraphics[width=4in]{Fig2.jpg} \\caption{L'amas de galaxies dit ``du boulet''. Il s'agit en fait de deux amas de galaxies en collision. \\`A la photo optique du syst\\`eme (qui montre les deux amas de galaxies comme une concentration de t\\^aches blanches) est superpos\\'ee en rouge l'\\'emission X du gaz chaud, et en bleu la masse totale projet\\'ee. La distribution de masse projet\\'ee sur le ciel correspond \\`a la masse des amas reconstruite par lentilles gravitationnelles. Le plus petit amas de galaxies \\`a droite semble avoir travers\\'e comme un boulet le gros amas \\`a gauche. La pointe de couleur rouge \\`a droite montre clairement une onde de choc en forme de c\\^one, qui permet de mesurer la vitesse supersonique de cette collision (ici 4700 km/s, ou Mach 3). Tous les composants de l'amas ne r\\'eagissent pas de la m\\^eme fa\\c{c}on dans la collision\\,: les galaxies et la mati\\`ere noire peuvent s'interp\\'en\\'etrer sans presque se voir. En revanche, le gaz chaud est frein\\'e, si bien que les deux composants gazeux des deux amas sont plus rapproch\\'es que les deux masses totales. Le comportement diff\\'erent dans la collision du gaz chaud, des galaxies et de la mati\\`ere noire permet de les s\\'eparer, et de tester les mod\\`eles. D'apr\\`es Clowe et al. (2006).} \\label{fig2}} \\end{figure} ",
        "conclusions": ""
    },
    "0910/0910.5497_arXiv.txt": {
        "abstract": "s{I summarize recent results from Smith, Hernandez-Monteagudo \\& Seljak\\,\\citep{Smithetal2009}, a study of the impact of nonlinear evolution of gravitational potentials in the LCDM model on the Integrated Sachs-Wolfe (ISW) contribution to the cross-power spectrum of the CMB and a set of biased tracers of the mass. We use a large ensemble of $N$-body simulations to directly follow the potentials and compare the results to analytic perturbation theory (PT) methods.  The PT predictions match our results to high precision for $k<0.2 \\kMpc$.  We analyze the CMB--density tracer cross-spectrum using simulations and renormalized bias PT, and find good agreement.  The usual assumption is that nonlinear evolution enhances the growth of structure and counteracts the linear ISW on small scales, leading to a change in sign of the CMB-LSS cross-spectrum at small scales. However, PT analysis suggests that this trend reverses at late times when the logarithmic growth rate $f=d\\ln D/d\\ln a<1/2$ or $\\Omega_m (z)<0.3$. Numerical results confirm these expectations and we find nonlinear enhancement of the ISW signal on small scales at late-times. On computing the total contribution to the angular spectrum, we find that nonlinearity and scale dependence of the bias are unable to influence the signal-to-noise of the current and future measurements.} ",
        "introduction": "\\label{sec:intro}  The temperature fluctuations in the Cosmic Microwave Background (CMB) are directly sensitive to the presence of Dark Energy, through the change in energy that a CMB photon experiences as it propagates through an inhomogeneous Universe with time evolving gravitational potentials, $\\pdot$. There are three main processes that give rise to changing potentials: the Integrated Sachs--Wolfe effect~\\citep[][ISW]{SachsWolfe1967}, which deals with the linear evolution of potentials; the Rees--Sciama effect~\\citep[][RS]{ReesSciama1968}, which is concerned with the nonlinear evolution of potentials; and the Birkinshaw--Gull effect~\\citep{BirkinshawGull1983}, where time evolving potentials arise due to mass flows. We define the ``nonlinear ISW effect'' as the sum of all three contributions.  Unfortunately, the nonlinear ISW effect can not be observed directly in the CMB auto-power spectrum, owing to the fact that it affects the low multipoles, where cosmic variance is large. It can however be directly observed by cross-correlating the CMB with tracers of the Large-Scale Structure (hereafter LSS). \\citep{CrittendenTurok1996}. This analysis has recently been performed by a number of groups using the WMAP data and several LSS measurements (e.g. SDSS, NVSS, 2MASS), with claims of up to $4\\sigma$ level detections \\citep{Giannantonioetal2008,Hoetal2008}. In a recent paper, \\citep{Granettetal2008b} measured the cross-correlation between super-structures and super-voids with the CMB. On stacking the signal they found a $\\sim4.5\\sigma$ detection, in multiple WMAP bands, and the sign of which appeared consistent with late-time ISW. This result is very puzzling when one considers signal-to-noise ($\\SN$) calculations within the LCDM model. These show that the maximum $\\SN\\sim7$ for full sky surveys \\citep{CrittendenTurok1996}. How then is it possible to obtain such high $\\SN$, given the partial sky coverage of current surveys?  One proposed solution is that nonlinear evolution of potentials and galaxy bias may increase the $\\SN$.  In this study we calculate the impact of nonlinear evolution of gravitational potentials on the angular cross-spectrum between a biased set of density tracers and the CMB. We ask whether such effects influence $\\SN$. We pursue a two-pronged attack: our first line of inquiry is analytic, and we use perturbation theory (PT) to predict $\\pdot$, and this we do in \\S\\ref{sec:PT}; our second line is to use a large ensemble of $N$-body simulations to directly follow $\\pdot$, and this we do in \\S\\ref{sec:nbody}. ",
        "conclusions": ""
    },
    "0910/0910.5174_arXiv.txt": {
        "abstract": "{ Hyper-velocity stars (HVS) are moving so fast that they are unbound to the Galaxy. Dynamical ejection by a supermassive black hole is favoured to explain their origin. } { Locating the place of birth of an individual HVS is of utmost importance to understanding the ejection mechanism. } { SDSS~J013655.91+242546.0 (J0136+2425 for short) was found amongst three high-velocity stars (drawn from a sample of more than 10 000 blue stars), for which proper motions were measured. A kinematical as well as a quantitative NLTE spectral analysis was performed. When combined with the radial velocity (RV) and the spectroscopic distance, the trajectory of the star in the Galactic potential was reconstructed. } { J0136+2425 is found to be an A-type main-sequence star travelling at $\\approx$590~\\kms, possibly unbound to the Galaxy and originating in the outer Galactic rim nowhere near the Galactic centre. } { J0136+2425 is the second HVS candidate with measured proper motion, besides the massive B star HD~271791, and also the second for which its proper motion excludes a Galactic centre origin and, hence, the SMBH slingshot mechanism. Most known HVS are late B-type stars of about 3 M$_\\odot$. With a mass of 2.45 M$_\\odot$, J0136+2425 resembles a typical HVS far more than HD~271791 does. Hence, this is the first time that a typical HVS is found not to originate in the Galactic centre. Its ejection velocity from the disk is so high (550~\\kms) that the extreme supernova binary scenario proposed for HD~271791 is very unlikely. } ",
        "introduction": "\\label{sec:intro} Stars travelling so fast that they escape from the Galaxy are an inevitable consequence of the presence of a supermassive black hole (SMBH) in a dense stellar environment \\citep{1988Natur.331..687H} such as the Galactic centre (GC). \\citet{1988Natur.331..687H} coined the term hyper-velocity star (HVS) for such an object. It took a long time until \\cite{2005ApJ...622L..33B}, \\cite{2005A&A...444L..61H} and \\cite{2005ApJ...634L.181E} discovered the first three HVSs serendipitously and the interest in these stars grew tremendously. Systematic searches for HVSs took advantage of the huge \\sdss database \\citep{2008ApJS..175..297A} for target selection and revealed a population of HVSs; the latest compilation lists 16 of these stars \\citep{2009ApJ...690.1639B}. These surveys are based on RV measurements alone. The known HVSs are non-uniformly distributed on the sky, as are their travel times. \\citet{2009ApJ...690L..69B} argue that this anisotropy supports their common origin being the Galactic centre, while \\citet{2009ApJ...691L..63A} point out that the overdensity of HVSs in the constellation Leo suggests that the anisotropic distribution and preferred travel time are the result of the tidal disruption of a dwarf galaxy in the Galactic potential.  Another issue is the distance of a HVS star, because blue horizontal branch (BHB) stars can not be easily distinguished from main-sequence (MS) stars for the HVSs in question since both types of stars populate the same region in the \\teff-$\\log{g}$-diagram \\citep{2008ASPC..392..167H}, but have different distances. It is generally assumed that the stars are main-sequence and not blue horizontal branch stars. Detailed spectroscopic analyses have confirmed these assumptions for HVS1, HVS3, HVS7, and HVS8 \\citep{2008A&A...480L..37P,2008A&A...488L..51P,2008ASPC..392..167H,2008ApJ...685L..47L}. In the absence of proper motion measurements, the trajectories have not been derived for any individual HVS so far. \\citet{2008A&A...483L..21H} succeeded in investigating the trajectory of the high-velocity B star HD~271791 from accurate proper motions, radial velocity, and spectroscopic distance. The star was found to be probably unbound to the Galaxy and to originate in the outer rim of the Galaxy rather than in its centre, demonstrating that a mechanism other than that of \\citet{1988Natur.331..687H} must operate. Hence, it is rewarding to measure proper motions of high-velocity stars. \\citet{2008ApJ...684.1143X} presented radial velocities for more than 10 000 blue stars from the \\sdss, which are a mix of blue horizontal branch, blue straggler, and main-sequence stars with effective temperature between roughly 7 000 and 10 000\\,K according to their colours. We focused on the 11 fastest stars in terms of large positive Galactic rest-frame (GRF) velocities to unravel their nature, distance, and kinematics from detailed quantitative spectral analyses and astrometry. Three stars exhibited significant proper motions\\footnote{Two of the stars are found to be a metal-poor straggler of population II and a spectroscopic binary, as described in a forthcoming paper}. Here we report that J0136+2425 is an A-type main sequence star travelling at $\\approx$590~\\kms, possibly unbound to the Galaxy and originating in the outer Galactic disk. ",
        "conclusions": "\\label{sec:conclusion}  We have reported the quantitative spectral analysis of a high-velocity star from the sample of faint blue stars in the halo of \\citet{2008ApJ...684.1143X}. The radial velocity, proper motion, and spectroscopic distance were derived and a detailed kinematical analysis was performed using the Galactic potential of \\citet{1991RMxAA..22..255A}.  J0136+2425 was found to be a rapidly rotating A star of solar composition and therefore classified as a main-sequence star of 2.45~M$_\\odot$. The kinematic analysis excludes an origin at the GC and hence the Hills mechanism for ejection of the star. Its place of origin was found to be in the outer part of the Galactic plane at Galactic radii between 12.5~kpc and 18~kpc. This is very similar to the case of the massive B star HD~271791 \\citep{2008A&A...483L..21H}, which was the first ultra-high-velocity star whose proper motion excludes a GC origin. The ejection mechanism for the star was proposed to be an extreme binary supernova, which also explained its enrichment in $\\alpha$ elements. However, we find no evidence of $\\alpha$-enhancement in J0136+2425 from the existing spectra. The star is more metal-rich than expected on average for an object born at such a large Galactocentric distance. Due to the presence of Galactic abundance gradients \\citep{2006ApJS..162..346R}, we expect the outer parts of the Galaxy to be less metal-rich than the Sun. The metallicity of J0136+2425 is compatible with the high end of the abundance distribution in the outer disk.  The required ejection velocity from the disk is so high (550~\\kms) that an extreme supernova binary scenario as proposed for HD~271791 \\citep{2008ApJ...684L.103P} is very unlikely.  Most known HVS are late B-type stars of about 3 M$_\\odot$. With a mass of 2.45 M$_\\odot$, J0136+2425 resembles such a typical HVS much more than HD~271791 ($M=$11 M$_\\odot$) does. Hence, this is the first time that a typical HVS is found not to originate in the GC and excludes the SMBH slingshot mechanism. Hence, typical HVS may have been ejected by different mechanisms other than that proposed by Hills.  Once more, this calls for an alternative ejection scenario to the Hills mechanism such as dynamical ejection from clusters or binary supernovae \\citep[see][]{2009MNRAS.396..570G}. Hence, we are left with the dynamical ejection scenario or tidal disruption of a satellite galaxy, as proposed by \\cite{2009ApJ...691L..63A}. They noticed a clustering of more than half of the known HVS in a region of 26$^\\circ$-diameter in the constellation of Leo ($l_{II}\\approx 230^\\circ$, $b_{II}\\approx 60^\\circ$). They suggested that if these star stems from a disruption event of a dwarf galaxy, more high velocity stars should be found in that area. However, J0136+2425 ($l_{II}=136^\\circ$, $b_{II}=-37^\\circ$) is located far from Leo and is therefore probably unrelated."
    },
    "0910/0910.0084_arXiv.txt": {
        "abstract": "The nova-like cataclysmic binary AE Aqr, which is currently understood to be a former supersoft X-ray binary and current magnetic propeller, was observed for over two binary orbits (78 ks) in 2005 August with the High-Energy Transmission Grating (HETG) onboard the {\\it Chandra X-ray Observatory\\/}. The long, uninterrupted \\Chandra\\ observation provides a wealth of details concerning the X-ray emission of AE Aqr, many of which are new and unique to the HETG. First, the X-ray spectrum is that of an optically thin multi-temperature thermal plasma; the X-ray emission lines are broad, with widths that increase with the line energy, from $\\sigma\\approx 1$ eV ($510~\\rm km~s^{-1}$) for \\ion{O}{8} to $\\sigma\\approx 5.5$ eV ($820~\\rm km~s^{-1}$) for \\ion{Si}{14}; the X-ray spectrum is reasonably well fit by a plasma model with a Gaussian emission measure distribution that peaks at $\\log T ({\\rm K})=7.16$, has a width $\\sigma=0.48$, an Fe abundance equal to $0.44$ times solar, and other metal (primarily Ne, Mg, and Si) abundances equal to $0.76$ times solar; and for a distance $d=100$ pc, the total emission measure $EM=8.0\\times 10^{53}~\\rm cm^{-3}$ and the 0.5--10 keV luminosity $L_{\\rm X}= 1.1\\times 10^{31}~\\rm erg~s^{-1}$. Second, based on the $f/(i+r)$ flux ratios of the forbidden ($f$), intercombination ($i$), and recombination ($r$) lines of the He $\\alpha $ triplets of \\ion{N}{6}, \\ion{O}{7}, and \\ion{Ne}{9} measured by Itoh et al.\\ in the {\\it XMM-Newton\\/} Reflection Grating Spectrometer spectrum and those of \\ion{O}{7}, \\ion{Ne}{9}, \\ion{Mg}{11}, and \\ion{Si}{13} in the \\Chandra\\ HETG spectrum, either the electron density of the plasma increases with temperature by over three orders of magnitude, from $n_{\\rm e}\\approx 6\\times 10^{10}~\\rm cm^{-3}$ for \\ion{N}{6} [$\\log T({\\rm K})\\approx 6$] to $n_{\\rm e}\\approx 1\\times 10^{14}~\\rm cm^{-3}$ for \\ion{Si}{13} [$\\log T ({\\rm K})\\approx 7$], and/or the plasma is significantly affected by photoexcitation. Third, the radial velocity of the X-ray emission lines varies on the white dwarf spin phase, with two oscillations per spin cycle and an amplitude $K\\approx 160~\\rm km~s^{-1}$. These results appear to be inconsistent with the recent models of Itoh et al., Ikhsanov, and Venter \\& Meintjes of an extended, low-density source of X-rays in AE Aqr, but instead support earlier models in which the dominant source of X-rays is of high density and/or in close proximity to the white dwarf. ",
        "introduction": "AE Aqr is a bright ($V\\approx 11$) nova-like cataclysmic binary consisting of a magnetic white dwarf primary and a K4--5 V secondary with a long 9.88 hr orbital period and the shortest known white dwarf spin period $P=33.08$ s \\citep{pat79}. Although originally classified and interpreted as a disk-accreting DQ Her star \\citep{pat94}, AE Aqr displays a number of unusual characteristics that are not naturally explained by this model. First, violent flaring activity is observed in the radio, optical, ultraviolet (UV), X-ray, and TeV $\\gamma$-rays. Second, the Balmer emission lines are single-peaked and produce Doppler tomograms that are not consistent with those of an accretion disk. Third, the white dwarf is spinning down at a rate $\\Pdot = 5.64 \\times 10^{-14}~{\\rm s~s^{-1}}$ \\citep{deJ94}. Although this corresponds to the small rate of change of $1.78~\\rm ns~yr^{-1}$, AE Aqr's spin-down is typically characterized as ``rapid'' because the characteristic time $P/\\Pdot\\approx 2\\times 10^7$ yr is short compared to the lifetime of the binary and because the spin-down luminosity $L_{\\rm sd}=-I\\Om\\Odot \\approx 1\\times 10^{34}~\\rm erg~s^{-1}$ (where $I\\approx 0.2\\Mwd\\Rwd^2 \\approx 2\\times 10^{50}~\\rm g~cm^2$ is the moment of inertia for a white dwarf of mass $\\Mwd =0.8~\\Msun$ and radius $\\Rwd = 7.0\\times 10^8$ cm, $\\Om = 2\\pi/P$, and $\\Odot = -2\\pi\\Pdot/P^2$) exceeds the secondary's thermonuclear luminosity by an order of magnitude and the accretion luminosity by two orders of magnitude.  Because of its unique properties and variable emission across the electromagnetic spectrum, AE Aqr has been the subject of numerous studies, including an intensive multiwavelength observing campaign in 1993 October [\\citealt{cas96}, and the series of papers in ASP Conf.\\ Ser.\\ 85 \\citep{buc94}]. Based on these studies, AE Aqr is now widely believed to be a former supersoft X-ray binary \\citep{sch02} and current {\\it magnetic propeller\\/} \\citep{wyn97}, with most of the mass lost by the secondary being flung out of the binary by the magnetic field of the rapidly rotating white dwarf. These models explain many of AE Aqr's unique characteristics, including the fast spin rate and rapid secular spin-down rate of the white dwarf, the anomalous spectral type of the secondary, the anomalous abundances \\citep{mau97}, the absence of signatures of an accretion disk \\citep{wel98}, the violent flaring activity \\citep{per03}, and the origin of the radio and TeV $\\gamma$-ray emission \\citep{kui97, mei03}.  To build on this observational and theoretical work, while taking advantage of a number of improvements in observing capabilities, during 2005 August 28--September 2, a group of professional and amateur astronomers conducted a campaign of multiwavelength (radio, optical, UV, X-ray, and TeV $\\gamma$-ray) observations of AE Aqr. Attention is restricted here to the results of the X-ray observations, obtained with the High-Energy Transmission Grating (HETG) and the Advanced CCD Imaging Spectrometer (ACIS) detector onboard the {\\it Chandra X-ray Observatory\\/}. \\citet{mau06} has previously provided an analysis and discussion of the timing properties of these data, showing that: (1) as in the optical and UV, the X-ray spin pulse follows the motion of the white dwarf around the binary center of mass and (2) during the decade 1995--2005, the white dwarf spun down at a rate that is slightly faster than predicted by the \\citet{deJ94} spin ephemeris. Here, we present a more complete analysis and discussion of the \\Chandra\\ data, providing results that in many ways reproduce the results of the previous {\\it Einstein\\/}, {\\it ROSAT\\/}, {\\it ASCA\\/}, and {\\it XMM-Newton\\/} \\citep{pat80, rei95, cla95, osb95, era99, cho99, ito06, cho06} and the subsequent {\\it Suzaku\\/} \\citep{ter08} observations of AE Aqr, but also that are new and unique to the HETG; namely, the detailed nature of the time-average X-ray spectrum, the plasma densities implied by the He $\\alpha $ triplet flux ratios, and the widths and radial velocities of the X-ray emission lines. As we will see, these results appear to be inconsistent with the recent models of \\citet{ito06}, \\citet{ikh06}, and \\citet{ven07} of an extended, low-density source of X-rays in AE Aqr, but instead support earlier models in which the dominant source of X-rays is of high density and/or in close proximity to the white dwarf.  The plan of this paper is as follows. In \\S 2 we discuss the observations and the analysis of the X-ray light curve (\\S 2.1), spin-phase light curve (\\S 2.2),  spectrum (\\S 2.3), and radial velocities (\\S 2.4). In \\S 3 we provide a summary of our results. In \\S 4 we discuss the results and explore white dwarf (\\S 4.1), accretion column (\\S 4.2), and magnetosphere (\\S 4.3) models of the source of the X-ray emission in AE Aqr. In \\S 5 we draw conclusions, discuss the bombardment model of the accretion flow of AE Aqr, and close with a few comments regarding future observations. The casual reader may wish to skip \\S 2.1--2.2, which are included for completeness, and concentrate on \\S 2.3--2.4, which contain the important observational and analysis aspects of this work. ",
        "conclusions": "Of the simple models considered above for the source of the X-ray emission of AE Aqr, the white dwarf and magnetosphere models are the most promising, while the accretion column model appears to be untenable: it fails to reproduce the \\Chandra\\ HETG X-ray light curves and radial velocities, and it requires an X-ray luminosity that is more than two orders of magnitude greater than observed in the 0.5--10 keV bandpass to heat the UV hot spots by reprocessing. A more intimate association between the X-ray and UV emission regions could resolve this problem, and it is interesting and perhaps important to note that such a situation is naturally produced by a bombardment and/or blobby accretion solution to the accretion flow. In particular, if the cyclotron-balanced bombardment solution applies to AE Aqr, it would naturally produce comparable X-ray and UV spot sizes and luminosities, comparable relative pulse amplitudes, comparable orbit-phase pulse time delays, the observed tight correlation between the X-ray and UV light curves \\citep{mau09}, and an accretion region that is heated to the observed $T\\sim 10^7$ K \\citep{woe92,woe93,woe96, fis01}. Such a solution to the accretion flow applies if (1) the specific accretion rate ($\\dot m =\\dot M/A$) onto the white dwarf is sufficiently low that the accreting plasma does not pass through a hydrodynamic shock, but instead is stopped by Coulomb interactions in the white dwarf atmosphere, and (2) the magnetic field of the white dwarf is sufficiently strong that the accretion luminosity can be radiated away by a combination of optically thin bremsstrahlung and line radiation in the X-ray waveband and optically thick cyclotron radiation in the infrared and optical wavebands. According to \\citet[equating $x_{\\rm s}$ from eqn.~17 with that from eqn.~20]{fis01}, the limit for the validity of cyclotron-balanced bombardment is $\\dot m\\la 1.19^{+0.69}_{-0.47}\\times 10^{-4}\\, B_7^{2.6}~\\rm g~cm^{-2}~s^{-1}$ for a white dwarf mass $\\Mwd=0.8\\pm 0.2\\, \\Msun $ and magnetic field strength $B_7=B_\\star/10^7$ G.  Are these conditions satisfied in AE Aqr? Perhaps. First, assuming that the observed X-ray luminosity is driven by accretion onto the white dwarf, the accretion rate $\\dot M=L_{\\rm X}\\Rwd/G\\Mwd \\approx 7.3\\times 10^{13}~\\rm g~s^{-1}$, hence the specific accretion rate $\\dot m\\approx 4.7\\times 10^{-5}\\, (f/0.25)^{-1}~\\rm g~cm^{-2}~s^{-1}$. Second, the strength of the magnetic field of the white dwarf in AE Aqr is uncertain, but based on the typical magnetic moments of intermediate polars ($10^{32}~\\rm G~cm^3$) and polars ($10^{34}~\\rm G~cm^3$), it should lie in the range $0.3~{\\rm MG}\\la B_\\star\\la 30~{\\rm MG}$. More specific estimates of this quantity include $B_\\star \\la 2$ MG, based on evolutionary considerations \\citep{mei02}; $B_\\star\\sim 1$--5 MG, based on the low levels of, and upper limits on, the circular polarization of AE Aqr \\citep[although these values are very likely lower limits, since these authors do not account explicitly for AE Aqr's low accretion luminosity and bright secondary]{cro86,sto92,bes95};  and $B_\\star \\sim 50$ MG, based on the (probably incorrect) assumption that the spin-down of the white dwarf is due to the pulsar mechanism \\citep{ikh98}. The above limit for the validity of cyclotron-balanced bombardment requires $B_\\star\\ga 7^{+1.5}_{-1.1}\\, (f/0.25)^{-0.35}$ MG, which is probably not inconsistent with the circular polarization measurements and upper limits. Conversely, the validity of the bombardment solution in AE Aqr could be tested with a more secure measurement of, or lower upper limit on, the magnetic field strength of its white dwarf.  Where do we go from here? Additional constraints on the global model of AE Aqr, and in particular on the connection between the various emission regions, can be supplied by simultaneous multiwavelength observations, including those obtained during our 2005 multiwavelength campaign, the analysis of which is at an early stage \\citep{mau09}. TeV observations with the Whipple Observatory 10-m telescope obtained over the epoch 1991--1995 \\citep{lan98} and the Major Atmospheric Gamma-Ray Imaging Cherenkov (MAGIC) 17-m telescope obtained during our campaign \\citep{sid08} provide only (low) upper limits on the TeV $\\gamma $-ray flux from AE Aqr, so it would be useful to unambiguously validate or invalidate earlier reports of detections of TeV flux from this source; in particular, a dedicated High Energy Stereoscopic System (HESS) campaign of observations, preferably in conjunction with simultaneous X-ray observations, is badly needed. Similarly, it would be useful to independently validate or invalidate the existence of the high-energy emission detected by \\citet{ter08} in the {\\it Suzaku\\/} X-ray spectrum of AE Aqr, and to establish unambiguously if it is thermal or nonthermal in nature. Going out on a limb, it is our bet that AE Aqr is not a TeV $\\gamma $-ray source and that the high-energy X-ray excess is thermal in nature, specifically that it is due to the stand-off shock present when the specific mass accretion rate onto the white dwarf is large (e.g., during flares) and the bombardment solution is no longer valid. Additional high resolution X-ray spectroscopic observations are also warranted, particularly to confirm or refute the high electron densities inferred from the flux ratios of the \\ion{Ne}{9}, \\ion{Mg}{11}, and \\ion{Si}{13} He $\\alpha $ triplets; it is, after all, {\\it a priori\\/} unlikely that each of these ratios lies on the knee of the theoretical $f/(i+r)$ flux ratio curves, and the trend of increasing density with increasing temperature is opposite to what is expected from a cooling flow (although it is consistent with the opposite, e.g., adiabatic expansion). It would be extremely helpful to apply the various Fe L-shell density diagnostics \\citep{mlf01,mlf03,mlf05} to AE Aqr, since they are typically less sensitive to photoexcitation, but the \\ion{Fe}{17} 17.10/17.05 line ratio for one is compromised by blending due to the the large line widths, and the other line ratios all require higher signal-to-noise ratio spectra than currently exists. It would be extremely useful to (1) determine the radial velocities of individual X-ray emission lines, rather than all the lines together; (2) resolve the individual lines on the white dwarf spin and orbit phases, (3) investigate the flare and quiescent spectra separately, (4) resolve individual flares, and (5) include the Fe K lines into this type of analysis. Such detailed investigations are beyond the capabilities of \\Chandra\\ or any other existing X-ray facility, but would be possible with future facilities such as {\\it IXO\\/} that supply an order-of-magnitude or more increase in effective area and spectral resolution."
    },
    "0910/0910.0751_arXiv.txt": {
        "abstract": "In this paper we analyze soft and hard X-ray emission of the 2002 September 20 M1.8 \\emph{GOES} class solar flare observed by \\emph{RHESSI} and \\textit{GOES} satellites. In this flare event, soft X-ray emission precedes the onset of the main bulk hard X-ray emission by $\\sim$5 min. This suggests that an additional heating mechanism may be at work at the early beginning of the flare. However \\emph{RHESSI} spectra indicate presence of the non-thermal electrons also before impulsive phase. So, we assumed that a dominant energy transport mechanism during rise phase of solar flares is electron beam-driven evaporation. We used non-thermal electron beams derived from \\emph{RHESSI }spectra as the heating source in a hydrodynamic model of the analyzed flare. We showed that energy delivered by non-thermal electron beams is sufficient to heat the flare loop to temperatures in which it emits soft X-ray closely following the \\emph{GOES} 1-8 \\AA ~light-curve. We also analyze the number of non-thermal electrons, the low energy cut-off, electron spectral indices and the changes of these parameters with time. ",
        "introduction": "It is commonly accepted that during the impulsive phase of the solar flare, non-thermal electron beams accelerated anywhere in the solar corona move along magnetic field lines to the chromosphere where they deposit their energy. Here, most non-thermal electrons lose their energy in Coulomb collisions while a small fraction ($\\sim 10^{-5}$) of the electrons' energy is converted into hard X-rays (HXR) by the bremsstrahlung process. The heated chromospheric plasma evaporates and radiates over a wide spectral range from hard X-rays or gamma rays to radio emission. These processes, called electron beam-driven evaporation model is believed to be a dominant energy transport mechanism during solar flares. One of the most spectacular and well observed manifestation of the described processes is soft X-ray (SXR) emission of magnetic loops observed e.g. by \\emph{Yohkoh} and \\emph{Hinode} satellites. Such a scenario implies that the hard and soft X-ray fluxes emitted by solar flares are generally related, as was first described by \\citet{Neup68}. Since HXR and microwave emissions are produced by non-thermal electrons and SXR are the thermal emission of a hot plasma, the Neupert effect confirms that non-thermal electrons deliver energy spent on plasma heating. Thus the HXR emission is directly related to the flux of the accelerated electrons whereas the SXR emission is related to the energy deposited by the non-thermal electron flux.  This interpretation, however, leads to a number of problems and questions concerning energy content and time relations between SXR and HXR fluxes (see e.g. \\citet{Den88},  \\citet{Den93}, and \\citet{McTier99}).  There are several papers that investigate temporal dependencies between the beginnings of SXR and HXR emissions. For example, \\citet{Mach86} and \\citet{Sch89} reported frequent strong SXR emission before the impulsive phase. Following these authors, on average, the SXR emission precedes the onset of HXR emission by about 2 min. \\citet{Ver02} analyzed 503 solar flares observed simultaneously in HXR, SXR and $H\\alpha$. In more than 90\\% of the analyzed flares, an increase of SXR emission began prior to the impulsive phase. On average the SXR emission starts 3 min before the hard X-ray emission.  An enhanced thermal emission preceding the onset of the hard X-rays may be indicative of a thermal preheating phase prior to the impulsive electron acceleration. Current-sheet models (e.g., \\citet{Heyv77}) of solar flares predict a preheating phase. \\citet{Li87} discussed the preheating phase of solar flares triggered by new emerging magnetic flux. They proposed four different types of reconnections to explain the preheating as well as impulsive phases of flares.   Another approach to this issue is multi-thread hydrodynamic modeling of solar flares. \\citet{Warr06} suggested that modeling a flare as a bunch of independently heated threads may simulate precedence of the SXR emission. The author successfully reproduced the temporal evolution of high temperature flare plasma in the Masuda flare of 1992 January 13.   An alternative mechanism to electron beam-driven evaporation, namely conducted driven evaporation, was developed lately by \\citet{Batt09}. They studied in detail the pre-flare phase of four solar flares using imaging and spectroscopy from the \\emph{RHESSI} satellite. These authors explain the time evolution of the observed emission for all analyzed events as an effect of saturated heat flux.   The tendency of the soft X-ray flux to appear before hard X-rays emission can be attributed also to the sensitivity threshold of the hard X-ray detectors \\citep{Den88}. At the beginning of the flare, the energy flux may be below the detection threshold of HXR emission.  In this paper we show, using unprecedented high sensitivity of the RHESSI detectors \\citep{Lin02} and a numerical model of flare, that early SXR emission observed prior to impulsive phase could be fully explained without any ad-hoc assumptions (at least for analyzed event). All necessary energy to explain the soft emission could be derived from observed HXR spectra. In Section 2 we describe the analyzed event. Section 3 presents the details of the HXR spectra fitting, numerical modeling and the results. The discussion and conclusions are presented in Section 5. ",
        "conclusions": "The start and quick increase of HXR $ \\ga 25\\rm\\, keV$ flux (defined as an impulsive phase of flare) is interpreted as an indication of the injection of the non-thermal electrons into the flaring loop and a beginning of the plasma heating by these electrons. However, often the SXR emission starts a few minutes earlier than the HXR which raises a question about the heating source in very early stage of the flare. Numerous solutions were proposed of which heating by thermal conduction and Multi-Stranded Loop model are the most well known. We showed that it is possible to fit soft (\\textit{GOES}) and hard (\\emph{RHESSI}) X-ray emissions of a solar flares well before beginning of the impulsive phase without any additional heating besides the heating by non-thermal electrons. This was made possible because of the unprecedented high sensitivity of the \\emph{RHESSI} detectors which are able to measure very low hard X-ray flux early in the flare. Part of the emission, mainly in the energy range $< 25 keV$, is non-thermal in nature and indicates the presence of non-thermal electrons. In this case of a M1.8 GOES class flare, the non-thermal electron energy fluxes of the order of $10^{26}~ergs/sec$, derived under the thick target interpretation, fully explains the required heating of the plasma and resulting increase in SXR emission.  The main limitations of our work are: crude estimation of the initial physical and geometrical parameters of the loop, errors in restoration of the RHESSI spectra and in GOES calibration, relative simplicity of the 1D hydro-dynamical model and single loop approximation of the event. While the synthesized GOES 1-8 \\AA ~light-curve closely follows the observed one, all limitations mentioned above caused small differences between observed and calculated fluxes in the 0.5-4 \\AA ~band.  Various authors considered classical heat conduction from the loop-top energy reservoir. \\citet{Batt09} fitted (\\emph{RHESSI}) spectra of four flares with a purely thermal emission and claim that at least for some flares electron beam heating does not work. In a case of flare analyzed by us purely thermal spectra seems to give too high temperatures ($> 20 ~MK)$ comparing to those obtained from \\emph{GOES} data. Additionally, even if the hard X-ray spectra observed with \\emph{RHESSI} provide no evidence for non-thermal particles this does not necessarily means that the electrons don't exists. As it was pointed e.g. by \\citet{BH09} the non-thermal hard X-ray emission associated with electron beam could be below \\emph{RHESSI's} level of detection. These authors estimated that in the case of analyzed event the non-thermal X-ray emission produced by an electron beam of sufficient energy flux to heat chromospheric plasma to the temperature and emission measure observed by \\emph{RHESSI} is below the observed background level. It is possible that for some flares heat conduction and for others beam driven evaporation works on the early phase of flare evolution. In our case the heating of the matter caused by non-thermal electrons is 2-4 orders higher then local conductive flux. \\citet{Sto08} presented analytic predictions of the X-ray $EM$ and $T_e$ expected in single and filamented flare loops for both mechanisms of evaporation and have tested these against real data consisting of 18 \\emph{RHESSI} microflares. Their results suggest beam heating in filamented loops to be in agreement with data. The peak temperatures were consistent with both single loop and multi-thread heating. The observed emission measures were mostly compatible with beam driving for a number of threads, but for 2 events EM were also compatible with single loop model.   Our results extend the standard model of SXR and HXR relationship to the early phases of solar flares and thus expands the number of flares consistent with the Neupert effect. These results also indicates that the process of electrons acceleration appears during the early stage of the flare, well before the impulsive phase. This was confirmed e.g. by \\citet{Asai06} who reported detection of the coronal non-thermal emission during the pre-impulsive phase of the X4.8 flare on 2002 July 23.  Our method of adjustment of the low energy cut-off $E_c$ in order to equalize synthesized and observed \\emph{GOES} fluxes in 1-8 \\AA ~channel can be considered as a new method of $E_c$ determination."
    },
    "0910/0910.2740_arXiv.txt": {
        "abstract": "We investigate the centrifugally driven curvature drift instability to study how field lines twist close to the light cylinder surface of an AGN, through which the free motion of AGN winds can be monitored. By studying the dynamics of the relativistic MHD flow close to the light cylinder surface, we derive and solve analytically the dispersion relation of the instability by applying a single particle approach based on the centrifugal acceleration. Considering the typical values of AGN winds, it is shown that the timescale of the curvature drift instability is far less than the accretion process timescale, indicating that the present instability is very efficient and might strongly influence processes in AGN plasmas. ",
        "introduction": "For studying AGN winds the fundamental problem relates to the understanding of a question: how the plasma goes through the Light Cylinder Surface (LCS), which is the hypothetical zone where the linear velocity of rotation equals the speed of light. This implies that the plasma particles, which move along quasi-straight magnetic field lines in the nearby area of the LCS, must reach the speed of light. Generally speaking no physical system can maintain such a motion and a certain twisting process of the magnetic field lines must operate on the LCS. On the other hand if the trajectories are given by the Archimedes spiral, then the particles can cross the LCS avoiding the light cylinder problem \\cite{r03}. An additional step in this investigation is to identify the appropriate mechanism that provides the twisting of the magnetic field lines, giving rise to the shape of the Archimedes spiral, and in turn insures that the dynamics is force-free.  Since the innermost region of AGNs rotates, the role of the Centrifugal Force (CF) appears interesting. The centrifugally driven outflows have been extensively studied. Generalizing the work developed in \\cite{gl97} it was shown that due to the centrifugal acceleration, electrons gain very high energies with Lorentz factors up to $\\gamma\\sim 10^8$ \\cite{osm7,ra8}. This implies that the energy budget in the AGN winds is very high.  The centrifugally driven parametric instability was first introduced in \\cite{incr1,incr3} for the Crab pulsar and AGN jets respectively. Another kind of the instability which might be induced by the CF is the so called Curvature Drift Instability (CDI). In \\cite{mnras} the two-component relativistic plasma has been considered to study the role of the centrifugal acceleration in the curvature drift instability for pulsar magnetospheres. The investigation demonstrated high efficiency of the CDI. To investigate the twisting process of magnetic field lines due to the CDI, we apply the method developed in \\cite{mnras,ff} to AGN winds.  The paper is arranged as follows. In Sect. 2, we introduce the curvature drift waves and derive the dispersion relation. In Sect. 3, the results for typical AGNs are presented and, in Sect. 4, we summarize our results. \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics[angle=0]{9710fig.ps}} \\caption{Two orthonormal bases are considered: i) cylindrical components of unit vectors, (${\\bf e}_\\Phi, {\\bf e}_R$, ${\\bf e}_x$); ii) unit vectors of the system rigidly fixed on each point of the curve, (${\\bf e}_r, {\\bf e}_{\\theta}$, ${\\bf e}_x$), respectively. $C$ is the center of the curvature. Hereafter, this set of coordinates is referred to as the field line coordinates.}\\label{fig} \\end{figure} ",
        "conclusions": "\\label{sec:summary}  We summarize the principal steps and conclusions of our study to be:  \\begin{enumerate}   \\item Considering the relativistic two-component plasma for AGN winds, the centrifugally driven curvature drift instability has been studied.   \\item Taking into account a quasi single approach for the particle dynamics, we linearized the Euler continuity and induction equations. The dispersion relation characterizing the parametric instability of the toroidal component of the magnetic field has been derived.  \\item By considering the proper frequency of the curvature drift modes, the corresponding expression of the instability increment has been obtained for the light cylinder region.  \\item The efficiency of the CDI has been investigated by adopting four physical parameters, namely: the wavelength, flow density and Lorentz factors of electrons, and the luminosity of AGNs.  \\item By considering the evolution process of accretion, the corresponding timescale has been estimated for a physically reasonable area in the parametric space $L-M_{BH}$. It was shown that the instability timescale was lower by many orders of magnitude than the evolution timescale, indicating extremely high efficiency of the CDI.  \\item Examining the instability from the point of view of the energy budget, we have seen that the sweepback of the magnetic field lines requires only a small fraction of the total energy, which means that the CDI is a realistic process.    \\end{enumerate}    \\begin{theacknowledgments} I thank professor G. Machabeli for valuable discussions. The research was supported by the Georgian National Science Foundation grant GNSF/ST06/4-096. \\end{theacknowledgments}"
    },
    "0910/0910.5483_arXiv.txt": {
        "abstract": "Many extensions of the standard model predict the existence of hidden sectors that may contain unbroken abelian gauge groups. We argue that in the presence of quantum decoherence photons may convert into hidden photons on sufficiently long time scales and show that this effect is strongly constrained by CMB and supernova data. In particular, Planck-scale suppressed decoherence scales $D \\propto \\omega^2/M_{\\rm Pl}$ (characteristic for non-critical string theories) are incompatible with the presence of even a single hidden U$(1)$. The corresponding bounds on the decoherence scale are four orders of magnitude stronger than analogous bounds derived from solar and reactor neutrino data and complement other bounds derived from atmospheric neutrino data. ",
        "introduction": "Physics beyond the standard model typically predicts the existence of hidden sectors containing extra matter with new gauge interactions. There is no compelling reason why these extra sectors should be very massive if their interaction with standard model matter is sufficiently weak. This can be accomplished, {\\it e.g.}, in extra dimensional extensions with a geometric separation of sectors or in models where tree-level interactions are absent and higher order corrections suppressed by the scale of messenger masses.  The presence of these light hidden sectors can have various observable effects.  For example, in the case of kinetic mixing~\\cite{Holdom:1985ag,Dienes:1996zr} between a hidden U$(1)$ and the electromagnetic U$(1)$, hidden sector matter can receive small fractional electro-magnetic charges. These {\\it mini-charged} particles can have a strong influence on early Universe physics and astrophysical environments~\\cite{Raffelt:1996wa}. If the hidden sector U$(1)$ is slightly broken by a Higgs or St\\\"uckelberg mechanism we can have a situation analogous to neutrino systems with characteristic oscillation patterns between photons and hidden photons over sufficiently long baselines~\\cite{Okun:1982xi}.  Here we concentrate on another effect induced by the presence of these light hidden sectors: the sensitivity of photon propagation to sources of quantum decoherence, {\\it e.g.}~quantum gravity effects~\\cite{Hawking:1982dj}. A heuristic picture describes space-time at the Planck scale as a foamy structure~\\cite{Wheeler:1957mu}, where virtual black holes pop in and out of existence on a time scale allowed by Heisenberg's uncertainty principle~\\cite{Hawking:1995ag}. This can lead to a loss of quantum information across their event horizons, providing an ``environment'' that might induce quantum decoherence of apparently isolated matter systems~\\cite{Hawking:1982dj}.  It is an open matter of debate whether quantum decoherence induced by a quantum theory of gravity would simultaneously preserve Poincare invariance and locality~\\cite{Hawking:1982dj,Hawking:1995ag,Ellis:1983jz,Banks:1983by,Unruh:1995gn}. A violation of energy and momentum conservation by particle reactions with a space-time foam could be reflected by an (energy dependent) effective refractive index in vacuo~\\cite{AmelinoCamelia:1996pj}. This could be tested, {\\it e.g.}, by the measurement of the arrival time of gamma rays or high energy neutrinos at different energies or by the propagation of ultra-high energy cosmic rays \\cite{GonzalezMestres:1997dr}.  If, on the other hand, Poincare invariance is preserved, the presence of a non-trivial space-time vacuum can still be signaled by decoherence effects in systems of stable elementary particles  \\cite{Ellis:1983jz}. A particularly interesting and well-studied case are neutrino systems, where the interplay between mixing, mass oscillation and decoherence can influence atmospheric, solar and reactor neutrino data \\cite{Lisi:2000zt,Gago:2002na,Fogli:2007tx}, as well as flavour composition of astrophysical high-energy neutrino fluxes~\\cite{Hooper:2004xr, Anchordoqui:2005gj}.  If hidden sectors contain unbroken abelian gauge groups it is also feasible that the system of the electromagnetic photon ($\\gamma_0$) and hidden photons ($\\lbrace\\gamma_i\\rbrace$ with $i\\geq1$) experience energy and momentum conserving decoherence effects:  transitions between different photon ``species'', $\\gamma_\\mu\\to\\gamma_\\nu$, are allowed by gauge and Poincare invariance. This is the only alternative system to neutrinos involving a stable and neutral elementary particle of the standard model, which can experience these decoherence effects on extremely long (cosmological) time scales.  The outline of this paper is as follows. We will start in Sect.~\\ref{sec:QD} with an outline of the Lindblad formalism of quantum decoherence. We will then discuss in Sect.~\\ref{sec:absorption} the effect of photon decoherence on the Planck spectrum of the cosmic microwave background and the luminosity distance of type Ia supernovae, respectively. This enables us to derive strong limits on various decoherence models. We comment in Sect.~\\ref{sec:propagation} on the interplay of decoherence and photon interactions and outline a possible mechanism to extend the survival probability of extra-galactic TeV gamma rays. We finally summarize in Sect.~\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  We have discussed the effects of quantum decoherence of photons in the presence of hidden sector U$(1)$s. Quantum decoherence in the system of photon and hidden photons could be induced by a foam-like structure of space-time in a quantum theory of gravity, where virtual black holes pop in and out of existence at scales allowed by Heisenberg's uncertainty principle.  We have shown that these decoherence effects are strongly constrained by the absence of photon disappearance in the cosmic microwave background.  Furthermore quantum decoherence as an additional source of starlight dimming can also be constrained by the luminosity distance measurements of type Ia supernovae. Consequently based on the standard $\\Lambda$CDM model we can derive constraints on the decoherence in the presence of hidden U$(1)$s.  In principle, the decoherence effect could be strong enough to influence the evaluation of cosmological data. However, color dependencies in the dimming via photon decoherence are not favored by the data and could be used to derive further constraints.  Our main results are summarized in  the upper panel of Fig.~\\ref{fig:exclusion} and compared with the present bounds from non observation of decoherence effects in neutrino systems. For models with decoherence parameter $D$ constant or decreasing with energy,  the bounds from from CMB distortions and SN dimming are the strongest. One must however bear in mind that while bounds from neutrino data are ``unavoidable'' once the quantum decoherence effects exist, the effects on photon propagation discussed here rely on the extra assumption of the presence of massless hidden U$(1)$s. We also note that, generically decoherence effects which grow with energy are better constrained from the higher energy atmospheric neutrino data.  We have also discussed the interplay between photon absorption and decoherence.  This effect can become important for distant TeV gamma ray point-sources if the decoherence length is much smaller than the photon interaction length. Assuming $N-1$ additional U$(1)$s this could leave characteristic features in gamma ray spectra in the form of step-like drops by factors $1/N$ or by an effective increase of the absorption length $\\lambda$ to $\\lambda_{\\rm eff} = N\\lambda$."
    },
    "0910/0910.2430_arXiv.txt": {
        "abstract": "We carried out an unbiased, spectroscopic survey using the low-resolution module of the infrared spectrograph (IRS) on board Spitzer targeting two 2.6 square arcminute regions in the GOODS-North field. IRS was used in spectral mapping mode with 5 hours of effective integration time per pixel. One region was covered between 14 and 21$\\mu m$ and the other between 20 and 35$\\mu m$. We extracted spectra for 45 sources. About 84\\% of the sources have reported detections by GOODS at $24\\mu m$, with a median $f_\\nu(24\\mu m)\\sim 100\\mu Jy$. All but one source are detected in all four IRAC bands, $3.6$ to 8 $\\mu$m. We use a new cross-correlation technique to measure redshifts and estimate IRS spectral types; this was successful for $\\sim 60\\%$ of the spectra. Fourteen sources show significant PAH emission, four mostly SiO absorption, eight present mixed spectral signatures (low PAH and/or SiO) and two show a single line in emission. For the remaining 17, no spectral features were detected. Redshifts range from $z\\sim 0.2$ to $z\\sim 2.2$, with a median of 1. IR Luminosities are roughly estimated from 24$\\mu m$ flux densities, and have median values of $2.2 \\times 10^{11} L_{\\odot}$ and $7.5 \\times 10^{11} L_{\\odot}$ at $z\\sim 1$ and $z\\sim 2$ respectively. This sample has fewer AGN than previous faint samples observed with IRS, which we attribute to the fainter luminosities reached here. ",
        "introduction": "The Spitzer Space Telescope \\citep{2004AdSpR..34..600W} has multiplied by several orders of magnitude the volumes previously surveyed in the  infrared for extragalactic objects \\citep{2008ARA&A..46..201S}.  Mid-infrared continuum surveys turned out to be especially fertile ground  because of the superb sensitivity and speed of the 24$\\mu m$ channel in the MIPS instrument \\citep{2004ApJS..154...25R}.  These mid-infrared surveys have established the dramatic evolution of galaxy populations over the last 10-12 billion years (e.g. \\citeauthor{2004ApJS..154..170L} \\citeyear{2004ApJS..154..170L}, \\citeyear{2005ApJ...632..169L} ; \\citeauthor{2006ApJ...637..727C} \\citeyear{2006ApJ...637..727C}; \\citeauthor{2007ApJ...668...45P} \\citeyear{2007ApJ...668...45P}) previously suggested by ISO surveys (e.g. \\citeauthor{1999A&A...351L..37E} \\citeyear{1999A&A...351L..37E}; \\citealt{2000MNRAS.316..749O}; \\citeauthor{2001A&A...372..364D} \\citeyear{2001A&A...372..364D}; \\citealt{2004MNRAS.351.1290R}), and yielded  rich samples of galaxies for followup.  Critical to the interpretation of the data is the availability of multiple-band detections, which are used to constrain the redshifts, luminosities and sometimes the powering mechanism of sources.  These constraints come about because the mid-infrared spectra of galaxies are often strongly structured (\\citeauthor{2007ApJ...656..770S} \\citeyear{2007ApJ...656..770S}; \\citeauthor{2007ApJ...656..148A} \\citeyear{2007ApJ...656..148A}), and sensitive to the presence of an AGN, to optical depth, and to heating and dust geometry.  By the same token, any single-band survey will have redshift-dependent  biases resulting in a sampling of the population that varies with redshift. The obvious way to avoid this sampling bias is to survey in the dispersed light, a technique pioneered from the ground using objective-prism imaging in the visible.  In this paper we report on a spectrally unbiased extragalactic census in the mid-infrared using the IRS instrument \\citep{2004ApJS..154...18H} on Spitzer, and covering two spectral windows, 14 to 21$\\mu m$ and 20 to 38$\\mu m$. Besides the specific results reported here, this survey illustrates the advantages of spectral mapping in the infrared. The placement of the windows allows us to attach an observed spectrum to the corresponding 24$\\mu m$ source, and could therefore potentially address the bolometric correction for each source.  The data  also allow us to use spectral shapes, features and emission lines to characterize all  detected sources, rather than only those sources  selected for follow-up by the inevitably rough criteria of filtering on colors within the mid-infrared or in combination with other bands (e.g. \\citeauthor{2004ApJS..154...75Y} \\citeyear{2004ApJS..154...75Y}; \\citeauthor{2006ApJ...651..101W} \\citeyear{2006ApJ...651..101W}; \\citeauthor{2007ApJ...671..323H} \\citeyear{2007ApJ...671..323H}; \\citeauthor{2008ApJ...677..957F} \\citeyear{2008ApJ...677..957F}). This unbiased examination of the emission line properties also  yields a fair assessment of the distribution of line strengths and line-to-continuum ratios.  Such an assessment constrains the frequency of line-dominated sources, and might even yield examples of sources radiating largely line emission, and therefore very rarely picked up in broad-band surveys. When our survey was proposed, a few examples of line-dominated sources in the mid-infrared were known, including NGC~1569 \\citep{2003ApJ...588..199L}.  Today the most striking case known is probably the intergalactic shock in Stephan's Quintet, which emits more than 20\\% of its total infrared emission in the main three molecular hydrogen pure rotational transitions \\citep{2006ApJ...639L..51A}. There is no reason to rule out scaled up versions of these systems, which might be detectable with Spitzer.  Source confusion is often  a limiting factor for infrared surveys, since the need to cool telescopes limits the size of the primary. The confusion limit appears in a variety of forms and definitions (\\citeauthor{1990LIACo..29..117H} \\citeyear{1990LIACo..29..117H}; \\citeauthor{2001ApJ...549..745R} \\citeyear{2001ApJ...549..745R}; \\citeauthor{2003MNRAS.338..555L} \\citeyear{2003MNRAS.338..555L}; \\citeauthor{2004ApJS..154...93D} \\citeyear{2004ApJS..154...93D}), but is ultimately dictated by the source density and its dependence on flux.  A spectrally dispersed survey offers the spectral dimension as a discriminator among closely spaced sources, since the strongly featured shapes of mid-infrared spectra would cause sources to peak almost always at different wavelengths. A method based on PSF fitting at different wavelengths allows to estimate accurately the confused sources positions at their peak wavelength as well as their contribution in flux to the neighboring sources. This spectral separation, solely relying on the knowledge of the PSF of the instrument, can be used as a prior to extract distinct spectra and thus disentangle sources, even if their broad-band images overlap substantially.  This paper describes the survey design, observations and data reduction in Sect. 2, describes the extraction of redshifts and spectral types from the data in Sect. 3, presents results on the distribution of redshifts and relation to broadband surveys in Sect. 4, and discusses those results in Sect. 5.  Throughout this paper, we assume a $\\Lambda$-CDM cosmology with $H_{0} = 70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{M} = 0.3$ and $\\Omega_{\\Lambda} = 0.7$. ",
        "conclusions": "We have obtained IRS spectra for 45 sources of very faint IR galaxies and determined redshifts for $\\sim 60\\%$ (28/45 sources) of them using a specifically designed cross-correlation method. We covered a domain of IR luminosities from about $7 \\times 10^{10} L_\\odot$ at $z = 0.5$ to about $10^{12} L_\\odot$ at $z \\sim 2.0$ poorly covered by IRS spectroscopy so far. At least 47\\% of our sample show starburst activity as PAH signatures (up to 21/45 sources) and 31\\% (14/45 sources) prove to be starburst dominated. The rest of the redshift identifications rest on silicate absorption feature (11\\% or 5/45 sources). We tentatively identify only 9\\% of our sample (4/45 sources) as AGN candidates. This small fraction of AGN in comparison to previous IRS spectroscopic surveys of high-redshifts galaxies is likely to originate from the fainter luminosities reached here. We find two unusual sources with only one prominent emission line detected over the spectral range covered. One of them present a strong [OIV] line (25.9$\\mu$m restframe) at $z = 0.3$ on top of a significant continuum while the other shows the [SIV] line (10.5$\\mu$m restframe) at $z = 2.08$ on a very weak continuum dominated by instrumental noise."
    },
    "0910/0910.1269_arXiv.txt": {
        "abstract": "The NEMO Collaboration installed and operated an underwater detector including prototypes of the critical elements of a possible underwater km$^3$ neutrino telescope: a four-floor tower (called Mini-Tower) and a Junction Box. The detector was developed to test some of the main systems of the km$^3$ detector, including the data transmission, the power distribution, the timing calibration and the acoustic positioning systems as well as to verify the capabilities of a single tridimensional detection structure to reconstruct muon tracks. We present results of the analysis of the data collected with the NEMO Mini-Tower. The position of photomultiplier tubes (PMTs) is determined through the acoustic position system. Signals detected with PMTs are used to reconstruct the tracks of atmospheric muons. The angular distribution of atmospheric muons was measured and results compared to Monte Carlo simulations. ",
        "introduction": "\\label{sec:introduction}   High energy neutrinos are considered optimal probes to identify the sources of high energy cosmic rays. Many indications suggest that cosmic objects, where acceleration of charged particles takes place, should also produce high energy gamma ray and neutrino fluxes. Indeed, $p \\gamma$ or $p p$ interactions responsible for $>$ TeV neutrinos and $\\gamma$ rays, are expected to occur in several astrophysical environments such as supernova remnants, microquasars, gamma-ray bursts and active galactic nuclei \\cite{ref:Aharonian_Science_2007}. For many sources the neutrino flux,  produced in the decay chains due to charge pions, is expected to be similar to the flux of high-energy gamma rays of hadronic origin, which can be measured by TeV gamma ray telescopes such as MAGIC \\cite{ref:MAGIC_2005}, HESS \\cite{ref:HESS_SNR_Nature2004} and VERITAS \\cite{ref:VERITAS_2008}. Because of the low expected neutrino fluxes from galactic and extragalactic sources \\cite{modelli}, the effective opening of the high energy neutrino astronomy era can really be made with detectors of km$^3$ scale. After the success of the first  generation of underwater/ice neutrino telescopes, such as BAIKAL \\cite{baikal} and AMANDA \\cite{amanda}, the construction of the first km$^3$ telescope, IceCube, \\cite{icecube} started at the South Pole. The detector is announced to be completed by 2011. In the Mediterranean Sea, the ANTARES telescope is taking data since 2006 in a partial configuration and since 2008 in its full set-up \\cite{antares}. The ANTARES Collaboration together with the NESTOR \\cite{nestor} and NEMO \\cite{nemo} Collaborations are conducting an intense R\\&D activity for the future km$^3$ Mediterranean telescope. Recently the three collaborations joined their efforts to design the KM3NeT undersea infrastructure that will host a km$^3$ telescope; its construction is expected to start by 2013 \\cite{km3net}.   The activity of the NEMO Collaboration was mainly focused on the search and characterization of an optimal site for the detector installation and on the development of key technologies for the km$^3$ underwater telescope to be installed in the  Mediterranean Sea. A deep sea site with optimal features in terms of depth and water optical properties was identified at a depth of 3500 m about 80 km off-shore from Capo Passero (Southern cape of Sicily).  A long term monitoring of the site has been carried out \\cite{sito}.  An other effort of the NEMO Collaboration was the definition of a feasibility study of the km$^3$ detector, which included the analysis of the construction and installation issues and the optimization of the detector geometry by means of numerical simulations.  To ensure an adequate process of validation of the proposed solutions, a technological demonstrator was built and installed off-shore the port of  Catania (Sicily, Italy). This project, called NEMO Phase-1, allowed the test and the qualification of the key technological elements (mechanics, electronics, data transmission, power distribution, acoustic positioning and time calibration system) proposed for the km$^3$ detector  \\cite{Migneco08}. Some of the technical solutions, further developed within the KM3NeT Consortium, have been proposed for the km$^3$ detector construction.   In this paper, after a detailed description of the NEMO Phase-1 lay-out and operation, we focus on the atmospheric muon data analysis procedure and present the main results. In particular, the atmospheric muon angular distribution was measured and compared with Monte Carlo simulations. The vertical muon flux was determined from the angular distribution, computing the related depth intensity relation. Results were compared with theoretical predictions and with results of other experiments. ",
        "conclusions": "The activities of the NEMO Collaboration have progressed with the achievement of major milestones: the realization and installation of the Phase-1 apparatus.  With this apparatus it was possible to test in deep sea the main technological solutions developed by the collaboration for the km$^3$ scale underwater neutrino telescope.  The angular distribution of atmospheric muons was measured and results were compared to Monte Carlo simulations. The vertical muon intensity was evaluated and compared with previous data and predictions, showing a good agreement."
    },
    "0910/0910.3746_arXiv.txt": {
        "abstract": "New study confirms conclusions made in \\cite{Esel2008}; according to it, there is a disturbed region expended along the CME propagation direction in front of a coronal mass ejection whose velocity $u$ is lower than the critical $u_C$ relative to the surrounding coronal plasma. The time difference brightness (plasma density) in the disturbed region smoothly decreases to larger distances in front of CME. A shock wave forms at u higher than $u_C$ in the front part of the disturbed region manifested as a discontinuity in radial distributions of the difference brightness. ",
        "introduction": "A coronal mass ejection (CME) structure in white light is often characterized by the following well-known features: a bright frontal structure (FS) that covers the region of decreased plasma density (cavity) that may includes a bright interior (core). However, besides the said features, another extended disturbed region defined by \\cite{Esel2007a} can exist immediately in front of a CME. The aim of our study is to investigate changes in the disturbed region form, when a CME velocity increases, and possibilities for formation of a shock wave in this case. ",
        "conclusions": "It has been shown that in front of a coronal mass ejection having a velocity $u$ lower than the critical $u_C$ relative to the surrounding coronal plasma there is a disturbed region expended along a direction of the CME propagation. The time difference brightness $\\Delta P$ in the disturbed region smoothly decreases up to larger distances in front of the CME. Given $u > u_C$, a discontinuity forms in distribution of difference brightness or plasma density in the disturbed region front part. Since the $u_C$ value is close to the local fast-mode MHD velocity, which in corona approximately equal to the Alfven one, the formation of such a discontinuity when $u_C$ is exceeded may be identified with the formation of a shock wave.      \\subsection*"
    },
    "0910/0910.4399_arXiv.txt": {
        "abstract": "{Previous analyses of lithium abundances in main sequence and red giant stars have revealed the action of mixing mechanisms other than convection in stellar interiors. Beryllium abundances in stars with Li abundance determinations can offer valuable complementary information on the nature of these mechanisms. } {Our aim is to derive Be abundances along the whole evolutionary sequence of an open cluster. We focus on the well-studied open cluster IC~4651. These Be abundances are used with previously determined Li abundances, in the same sample stars, to investigate the mixing mechanisms in a range of stellar masses and evolutionary stages.} {Atmospheric parameters were adopted from a previous abundance analysis by the same authors. New Be abundances have been determined from high-resolution, high signal-to-noise UVES spectra using spectrum synthesis and model atmospheres. The careful synthetic modeling of the Be lines region is used to calculate reliable abundances in rapidly rotating stars. The observed behavior of Be and Li is compared to theoretical predictions from stellar models including rotation-induced mixing, internal gravity waves, atomic diffusion, and thermohaline mixing.} {Beryllium is detected in all the main sequence and turn-off sample stars, both slow- and fast-rotating stars, including the Li-dip stars, but is not detected in the red giants. Confirming previous results, we find that the Li dip is also a Be dip, although the depletion of Be is more modest than for Li in the corresponding effective temperature range. For post-main-sequence stars, the Be dilution starts earlier within the Hertzsprung gap than expected from classical predictions, as does the Li dilution. A clear dispersion in the Be abundances is also observed. Theoretical stellar models including the hydrodynamical transport processes mentioned above are able to reproduce all the observed features well. These results show a good theoretical understanding of the Li and Be behavior along the color-magnitude diagram of this intermediate-age cluster for stars more massive than 1.2 M$_{\\odot}$. } {} ",
        "introduction": "\\label{sec:int}  The classical theory of stellar evolution, which only allows for mixing in stellar convective layers, fails to explain nearly all Li abundance patterns observed so far at the surface of low-mass stars. Indeed, in contradiction to the classical predictions, the vast majority of F- and early G-type stars (including the Sun) deplete, often quite severely, their surface Li abundance during their main-sequence lifetime, revealing transport processes beyond convection. A-type stars are also concerned, although in these objects the signatures of non-standard Li depletion, occurring in their interior during the main sequence, appear at their surface only when they cross the Hertzsprung gap and become subgiants.  Lithium depletion is seen both in field \\citep[e.g.,][and references therein]{Herbig65,Duncan81, PLP94,Leb99,Chen01,LRe04} and in open cluster stars \\citep[e.g.][and references therein]{Wal65,Zappala72,Cayrel84,Soderblom90, Balachandran95,Jones97,PRP01,SR05}. As cluster stars have well-defined masses and, most important, share the same age and initial chemical composition, they are ideal targets for investigating mixing processes as a function of stellar mass and evolutionary status. They have therefore received considerable attention in the literature. Observations of the Li evolution along the color magnitude diagram (CMD) of open clusters have revealed a number of interesting features, such as the Li-dip \\citep{Wal65,BT86a}, the lack of abundance trends with age among A-type stars \\citep{BurCou98,BurCou00}, the presence of a spread of Li among solar-type stars in relatively old clusters like M67 \\citep{PRP97,JFS99}, and the Li variations with mass and evolutionary status as in the case of IC 4651 \\citep{PRZ04}.  Many different physical mechanisms have been proposed to explain these observations: atomic diffusion \\citep{Mic86,Mi04}, mass loss \\citep{Hobbs89,SwFa92}, rotation-induced mixing \\citep{CVZ92,DP97,TC98,PTCF03,TV03}, mixing by internal gravity waves \\citep{GLS91,MoSc00,Young03}, or a combination of some of these processes. So far, only the hydrodynamic stellar models that combine the effects of meridional circulation, shear turbulence, internal gravity waves, and atomic diffusion have been able to account for Li observations over a broad range of stellar masses, ages, and evolutionary status \\citep{TC03,TC05,ChT05,ipCT08}. These models, which use a prescription for the wave excitation that reproduces the solar-p modes \\citep{Goldreich94}, explain other observational constraints concomitantly, like the behavior of C, N, O, and Be across the Li dip in young clusters, such as the Hyades, or the internal rotation profile of the Sun \\citep{ChT05}.  Additional clues can help ascertain the nature and description of transport mechanisms of chemicals and angular momentum inside low-mass stars. In particular the simultaneous determination of Li and Be abundances in the same stars strengthens the constraints by performing a stellar tomography. Such a powerful diagnosis is possible because Li and Be burn at different temperatures ($\\sim$ 2.5$\\times 10^{6}$ K for Li and $\\sim$ 3.5$\\times 10^{6}$ K for Be), which correspond to different depths in the stellar interior \\citep[see e.g.][]{ipDPC00}.  \\begin{figure} \\centering \\includegraphics[width=7cm]{12957f01.eps} \\caption{Color magnitude diagram of IC 4651. Sample stars are shown as filled circles. The $ubvy$ photometry is from \\citet{Mei00} and obtained by means of the WEBDA database.} \\label{fig:cmd} \\end{figure}  Although the simultaneous study of Li and Be in the same stars is a powerful tool, Be has been determined in far fewer stars than Li has. In particular, despite the advantages of cluster stars mentioned above, Be has been investigated in only a few open clusters \\citep[e.g.][and references therein]{Boesgaard77,BAK04, GL95,Ran02,Ran07}, while field stars have been more extensively studied \\citep[e.g.][and references therein]{Boesgaard76b,SBKD97,San04b,BoKr09}.  Boesgaard and collaborators discovered a Be dip and a correlation between Li and Be depletion in mid F-stars in the field \\citep{DBS98} and also in the Hyades, Coma, and Praesepe clusters \\citep{BK02,BAK04}. They also show that for cooler G stars there is little or no Be depletion, even in Li-depleted objects up to $\\sim$ 5500 K \\citep[and references therein]{BAKDS04}. For even cooler stars, below $\\sim$ 5500 K, \\citet{San04b} show a Be decline correlated with decreasing temperature. These findings were confirmed by \\citet{Ran02,Ran07}, who investigated Li and Be in late F- and early G-type stars in some open clusters over a relatively broad age range (IC 2391 -- 50 Myr, NGC 2516 -- 150 Myr, Hyades -- 600 Myr, IC 4651 -- 1.7 Gyr, and M67 -- 4.5 Gyr).  \\begin{table} \\caption{Observational data of the sample stars.} \\label{tab:log} \\centering \\begin{tabular}{cccccc} \\noalign{\\smallskip} \\hline\\hline \\noalign{\\smallskip} Star & $v$  & ($b-y$) & obs. date & t$_{\\rm exp}$ (s) & S/N \\\\ \\hline E3    & 12.05 & 0.38 & 10.may.2001 & 2 $\\times$ 2100 & 90 \\\\ E5    & 12.82 & 0.33 & 09.may.2001 & 2 $\\times$ 3000 & 90 \\\\ E7    & 14.17 & 0.40 & 26.mar.2000 & 4 $\\times$ 3600 & 60 \\\\ E14   & 13.12 & 0.34 & 09.may.2001 & 2 $\\times$ 3300 & 95 \\\\ E15   & 13.52 & 0.38 & 10.apr.2001 & 2 $\\times$ 4500 & 85 \\\\ E19   & 12.24 & 0.42 & 08.apr.2001 & 2 $\\times$ 2100 & 85 \\\\ E34   & 13.42 & 0.36 & 11.may.2001 & 2 $\\times$ 4500 & 90 \\\\ E45   & 14.18 & 0.44 & 29.mar.2000 & 3 $\\times$ 3600 & 40 \\\\ E56   & 12.18 & 0.38 & 08.jun.2001 & 2 $\\times$ 2100 & 85 \\\\ E64   & 13.72 & 0.37 & 06.jun.2001 & 2 $\\times$ 4500 & 100 \\\\ E79   & 13.55 & 0.37 & 09.apr.2001 & 2 $\\times$ 4500 & 90 \\\\ E86   & 13.74 & 0.39 & 06.jun.2001 & 2 $\\times$ 4500 & 100 \\\\ E95   & 11.95 & 0.52 & 08.jun.2001 & 2 $\\times$ 2100 & 80 \\\\ E99   & 12.38 & 0.34 & 08.jun.2001 & 2 $\\times$ 2100 & 90 \\\\ T1228 & 12.94 & 0.35 & 08.jun.2001 & 2 $\\times$ 3300 & 90 \\\\ E25   & 12.65 & 0.37 & 07.jun.2001 & 2 $\\times$ 2100 & 85 \\\\ T2105 & 13.95 & 0.43 & \\multicolumn{3}{c}{\\citet{Ran02}}  \\\\ \\hline E8    & 10.67 & 0.66 & 10.may.2001 & 1500 & 50 \\\\ E12   & 10.32 & 0.64 & 15.apr.2001 & 2 $\\times$ 1200 & 75 \\\\ E60   & 10.87 & 0.67 & 09.apr.2001 & 2 $\\times$ 1800 & 70 \\\\ E98   & 10.88 & 0.66 & 11.may.2001 & 2 $\\times$ 1800 & 65 \\\\ T812  & 11.06 & 0.67 & 07.jun.2001 & 2 $\\times$ 1800 & 70 \\\\ \\noalign{\\smallskip} \\hline \\multicolumn{6}{l}{The Str\\\"omgren photometry is from \\citet{Mei00}. The S/N, per } \\\\ \\multicolumn{6}{l}{resolution element (4 pixels), was measured in the Be region. The} \\\\ \\multicolumn{6}{l}{ five stars grouped at the end of this and the following tables are the} \\\\ \\multicolumn{6}{l}{ red giants.} \\\\ \\end{tabular} \\end{table}   To improve our understanding of the mixing processes, adding new Be data covering a large fraction of the color magnitude diagram of well-studied open clusters is important. The ambitious aim is to compare the observations with models for stars over as wide as possible a range of mass and evolutionary status. In this context, the open cluster IC 4651 is a very good test case. It has well-determined membership \\citep{Mei02}, and precisely known metallicity and chemical composition \\citep[{[Fe/H]} = +0.11;][]{PRZ04,CBGT04,PP08}. Age determinations for this cluster vary between 1.2$\\pm$0.2 Gyr \\citep{Biaz07} and 1.7$\\pm$0.15 Gyr \\citep{Mei02}. A color magnitude diagram of IC 4651 is shown in Fig. \\ref{fig:cmd}.  Important is that the Li evolution along the CMD of this cluster has already been investigated in detail by \\citet{PRZ04}. Well-determined patterns have been found along its evolutionary sequence \\citep[see Fig.~3 by][]{PRZ04}. In a sequence of increasing stellar mass, the cluster indeed has an Li ``plateau'' around solar mass stars, followed by a well-determined Li dip, after which the maximum Li content is observed, at meteoritic value. A sudden Li drop follows along the Hertzsprung gap, possibly indicating the onset of dilution (the so-called first dredge-up). Some of the giants show noticeable differences in Li abundance while sharing other stellar characteristics. It has been shown that models of rotating stars by \\citet{CT99} and \\citet{PTCF03} were able to reproduce the Li behavior both of main sequence stars lying on the blue side of the dip and of evolved stars.  \\begin{table*} \\caption{Stellar parameters and abundances of Li and Be} \\label{tab:par} \\centering \\begin{tabular}{cccccccc} \\noalign{\\smallskip} \\hline\\hline \\noalign{\\smallskip} Star &  T$_{\\rm{eff}}$ & Spec.   & log g & $\\xi$ & $v~sini$ & A(Be) & A(Li) \\\\ &   (K)         & or Phot. & (dex) & (km s$^{-1}$)  & (km s$^{-1}$) & (dex) &  (dex)    \\\\ \\hline E3  & 6550 & Phot. & 3.90 & 2.10 & 29.9 & 1.21  & 1.64 \\\\ E5  & 6930 & Phot. & 4.20 & 2.00 & 23.9 & 1.36  & 3.47 \\\\ E7  & 6300 & Spec. & 4.30 & 1.10 &  2.1 & 1.11  & 2.83 \\\\ E14 & 6860 & Phot. & 4.20 & 1.90 & 34.1 & 0.76  & $\\leq$2.07 \\\\ E15 & 6540 & Phot. & 4.20 & 1.70 & 10.0 & 0.64  & $\\leq$1.93 \\\\ E19 & 6280 & Phot. & 3.90 & 2.10 & 32.1 & 0.41  & $\\leq$1.66 \\\\ E34 & 6640 & Phot. & 4.20 & 1.90 & 24.2 & 0.11  & $\\leq$1.92 \\\\ E45 & 6350 & Spec. & 4.30 & 1.10 &  4.2 & 1.08  & 2.80 \\\\ E56 & 6520 & Phot. & 3.90 & 2.10 & 27.8 & 1.21  & $\\leq$2.16 \\\\ E64 & 6650 & Spec. & 4.30 & 1.70 &  9.2 & 0.76  & 2.24  \\\\ E79 & 6620 & Spec. & 4.30 & 1.70 & 21.8 & 0.51  & $\\leq$1.91 \\\\ E86 & 6600 & Spec. & 4.30 & 1.70 & 14.0 & 0.96  & 2.20  \\\\ E95 & 5800 & Spec. & 3.50 & 1.70 & 12.0 & 0.61  & 2.19  \\\\ E99 & 6830 & Phot. & 4.00 & 2.00 & 28.1 & 1.21  & $\\leq$2.38 \\\\ T1228 & 6770 & Phot. & 4.20 & 1.70 & 5.7 & 1.21 & 3.18  \\\\ E25 & 6680 & Phot. & 4.00 & 2.00 & 21.3 & 1.11 &  3.29  \\\\ T2105 & 6110 & Phot. & 4.44 & 1.10 & 8.0 & 1.11 & 2.82 \\\\ \\hline E8  &  4900  & Spec. & 2.70 & 1.30 & 0.1 & -- & $\\leq$0.09 \\\\ E12 &  5000  & Spec. & 2.70 & 1.50 & 1.1 & $\\leq$$-$0.70 & $\\leq$0.35 \\\\ E60 &  4900  & Spec. & 2.90 & 1.40 & 1.9 & $\\leq$$-$0.70 & 0.42 \\\\ E98 &  4900  & Spec. & 3.00 & 1.40 & 1.0 & $\\leq$$-$0.50 & 0.72 \\\\ T812 &  5000 & Spec. & 3.00 & 1.60 & 0.5 & $\\leq$$-$0.70 & $\\leq$0.38 \\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{table*}  In this work we present Be abundances for the same stars analyzed by \\citet{PRZ04}. This is the first time Be abundances have been determined for stars along the whole evolutionary sequence of an open cluster, including solar-type, Li-dip, turn-off, subgiant, and red giant stars. With these new results, the mixing processes in different stellar masses can be investigated in greater detail. ",
        "conclusions": "\\label{sec:con}  Beryllium abundances were calculated for twenty-one main sequence, turn-off, subgiant, and giant stars belonging to the open cluster IC 4651. This is the first time that such an analysis has been carried out along the whole evolutionary sequence of an open cluster.  The Be abundances are found to closely follow the behavior of the Li abundances as determined by \\citet{PRZ04}. The coolest main sequence stars observed present no Be abundance dispersion, and a well-defined Be dip overlaps the well-documented Li dip for F-type stars. The Be abundances of post-main-sequence stars confirm that first-dredge up dilution starts earlier than the expected classically and present a significant dispersion. This is consistent with the Li behavior in evolved stars in IC~4651, in NGC~3680, and in the field.  Our observations are compared with theoretical predictions. For the main sequence and turnoff stars lying on the hot side of the Li-dip, as well as for the evolved stars, we compared the observations with results of new hydrodynamical stellar models calculated with STAREVOL V3.1 (Charbonnel \\& Lagarde, in preparation). These models include the effects of atomic diffusion, meridional circulation, shear turbulence, and thermohaline mixing. (Wave-induced transport is negligible in the case of these stars that have very narrow convective envelope while on the main sequence.) For the main sequence stars lying on the cool side of the Li dip, we compared the observations with the 1.2 M$_{\\odot}$ model of \\citet{ChT05} which also includes the transport of angular momentum by internal gravity waves.  The models reproduce all the Li and Be features very nicely along the CMD of IC~4651. The dispersion of Li and Be abundances on the blue side of the dip and in evolved stars is very well explained when accounting for a dispersion in the initial values of the stellar rotational velocity as observed in young open clusters. On the other hand, the lack of abundance dispersion in lower mass stars is expected to be caused by the impact of internal gravity waves.  The success in explaining the Li and Be abundances along the whole evolutionary sequence of IC~4651 is very encouraging. It shows that important steps have been taken towards the proper understanding of the physical mechanisms acting during the stellar evolution."
    },
    "0910/0910.5740_arXiv.txt": {
        "abstract": "Current models of pulsar gamma-ray emission use the magnetic field of a rotating dipole in vacuum as a first approximation to the shape of plasma-filled pulsar magnetosphere. In this paper we revisit the question of gamma-ray light-curve formation in pulsars in order to ascertain the robustness of the ``two-pole caustic\" and ``outer gap\" models based on the vacuum magnetic field. We point out an inconsistency in the literature on the use of the relativistic aberration formula, where in several works the shape of the vacuum field was treated as known in the instantaneous corotating frame, rather than in the laboratory frame. With the corrected formula, we find that the peaks in the light curves predicted from the two-pole caustic model using the vacuum field are less sharp. The sharpness of the peaks in the outer gap model is less affected by this change, but the range of magnetic inclination angles and viewing geometries resulting in double-peaked light curves is reduced. In a realistic magnetosphere, the modification of field structure near the light cylinder due to plasma effects may change the shape of the polar cap and the location of the emission zones. We study the sensitivity of the light curves to different shapes of the polar cap for static and retarded vacuum dipole fields. In particular, we consider polar caps traced by the last open field lines and compare them to circular polar caps. We find that the two-pole caustic model is very sensitive to the shape of the polar cap, and a circular polar cap can lead to four peaks of emission. The outer-gap model is less affected by different polar cap shapes, but is subject to big uncertainties of applying the vacuum field near the light cylinder. We conclude that deviations from vacuum field can lead to large uncertainties in pulse shapes, and a more realistic force-free field should be applied to the study of pulsar high energy emission. ",
        "introduction": "Pulsars are rotating neutron stars (NS) with very strong magnetic fields ($B\\sim10^9-10^{12}$G). Nearly two thousand pulsars are known, mostly in the radio band. Compton Gamma Ray Observatory (CGRO) firmly detected seven gamma-ray pulsars, and another three with less confidence [see \\citet{thom04} for a review]. This number is rapidly increasing with the start of operations of Fermi Gamma-ray Space Telescope (e.g., \\citealp{fermi08, FermiRelease09a, FermiRelease09b, FermiCatalog}). Gamma-ray light curves of pulsars are typically double-peaked, generally out of phase with the radio, and have substantial off-peak emission. Several theoretical models have been developed to explain the nature of this emission, namely, the polar-cap model \\citep{rs75,hte78,dh82,dh96}, the slot-gap model (SG, or two-pole caustic, TPC for short; \\citealp{as79,arons83,mh03,mh04,dr03,dhr04}), and the outer-gap (OG) model \\citep{chr86a,chr86b,ry95,yadi97,crz00}. In these models, particles are accelerated in the ``gap\" regions, where strong electric fields are developed due to a deficit of charge. Gamma-ray emission is interpreted as the curvature and/or inverse Compton radiation from ultra-relativistic particles accelerated in these gaps. The models differ in the location of the gaps in the magnetosphere. The polar cap model has narrow beam size and has difficulty in producing extended light curves, whereas SG/TPC and OG models can reasonably well reproduce double-peaked, extended light curves. Recent works have shown that OG model can also reproduce the high-energy spectrum for Crab and Vela pulsars \\citep{hirotani07,tc07,tcc07,tcs08}.  Special relativistic effects, such as the aberration of photon emission direction and photon travel time delay, are important for calculations of gamma-ray light curves when the emission zone extends far from the NS surface. In the TPC model, the emission zone is assumed to be along the last open field lines (LOFLs), extending from the polar cap to some cut-off radius. In the OG model, the emission zone is along the open field lines and extends from the null charge surface to the light cylinder\\footnote{Recent version of OG model allows inner boundary of emission zone to shift toward the NS surface [e.g. \\citet{tcs08}].} (LC, $R_{\\rm LC}=c/\\Omega$). In both models, the relativistic effects are essential to forming ``caustics\" in the sky map, which appear as sharp peaks in the light curve. The caustics arise when photons emitted from different regions of the magnetosphere happen to arrive to the observer at the same time. The presence and appearance of caustics is sensitive to both the relativistic effects and the geometry of the emission zone.  All existing calculations of pulsar gamma-ray light curves assume that pulsar magnetospheric geometry can be represented by a vacuum dipole magnetic field. However, the pulsar magnetosphere is filled with plasma \\citep{gj69}, and plasma currents should modify the field structure. The magnetosphere should then consist of the open and closed field line regions, separated by thin current sheets. In contrast, all field lines of the vacuum field, including those that travel beyond the LC, are formally closed, and no current sheets exist. Numerical solutions of the structure of plasma-filled magnetosphere are now known in the limit of force-free (FF) MHD for axisymmetric rotators \\citep{ckf99, Gruzinov05, Timokhin06, McKinney06}, and, recently, for three-dimensional oblique rotators as well \\citep{as06, cont09}. The FF field clearly demonstrates the current sheet structure, and its geometry differs substantially from the vacuum field near the LC [see \\citet{as08} for a review]. Hence, the more realistic FF field geometry can lead to modifications of gamma-ray light curves [see \\citet{bs09a} for preliminary results]. On the other hand, the vacuum field has been used over the years to obtain light curves that compare very favorably to the existing data. It can be argued that if the emission comes from the regions in the magnetosphere that are not too close to the light cylinder, the field geometry there may be well approximated by the vacuum field. This raises the question of how reliable and robust are the vacuum field light curves, and whether perturbations introduced by the presence of plasma in the magnetosphere would strongly affect the result.  Although our ultimate goal is to study gamma-ray emission using the force-free field, in this paper we concentrate on the modeling of light curves using vacuum magnetic field only. The results with the FF field will be presented in the companion paper \\citep{bs09b}. We feel it is necessary to clarify a number of points before moving forward. As we try to reproduce the sky maps and light curves using vacuum field geometry, we find that there are ambiguities in the literature on the use of the aberration formula. Since the aberration effect is crucial to the formation of caustics in any field geometry, in this paper we clarify the applicability of aberration formulas and compare their influences on the sky maps and light curves. In order to investigate the potential effects of plasma on the formation of light curves, we also study the sensitivity of the vacuum light curves to variations in the shape of the polar cap. Even if the field geometry in the bulk of the magnetosphere could be approximated by the vacuum field, it is the behavior of the field lines near the light cylinder that determines the shape of the polar cap and thus the location of the magnetospheric emission zones. The plasma effects near the light cylinder can then undermine light curve modeling that uses the polar caps of the vacuum field. In this paper we show that the appearance of the sky map and light curves is indeed very sensitive to both the field geometry and the geometry of the emission zones, which suggests that revisiting theoretical models with the more realistic FF field is essential.  This paper is structured as follows. We begin with the vacuum magnetic field formulas for pulsar magnetosphere in section 2, and then discuss the effect of aberration in section 3. In section 4, we construct the shape of the polar caps. We present comparisons of sky maps and light curves between different aberration formulas as well as different polar cap shapes for the two-pole caustic and the outer-gap models in section 5. In section 6 we summarize our results. ",
        "conclusions": ""
    },
    "0910/0910.0954_arXiv.txt": {
        "abstract": "Hard X-ray surveys have  proven remarkably efficient in detecting intermediate polars and asynchronous polars, two of the rarest type of cataclysmic variable (CV). Here we present a global study of hard X-ray selected intermediate polars and asynchronous polars, focusing particularly on the  link between hard X-ray properties and spin/orbital periods. To this end, we first construct a new sample of these objects by cross-correlating candidate sources detected in \\textit{INTEGRAL}/IBIS observations against catalogues of known CVs. We find 23 cataclysmic variable matches, and also present an additional 9 (of which 3 are definite) likely magnetic cataclysmic variables (mCVs) identified by others through optical follow-ups of IBIS detections. We also include in our analysis hard X-ray observations from \\textit{Swift}/BAT and \\textit{SUZAKU}/HXD in order to make our study more complete. We find that most hard X-ray detected mCVs have $P_{spin}/P_{orb}<0.1$ above the period gap. In this respect we also point out the very low number of detected systems in any band between $P_{spin}/P_{orb}=0.3$ and $P_{spin}/P_{orb}=1$ and the apparent peak of the $P_{spin}/P_{orb}$ distribution at about $0.1$. The observational features of the $P_{spin} - P_{orb}$ plane are discussed in the context of mCV evolution scenarios. We also present for the first time evidence for correlations between hard X-ray spectral hardness and $P_{spin}$, $P_{orb}$ and $P_{spin}/P_{orb}$. An attempt to explain the observed correlations is made in the context of mCV evolution and accretion footprint geometries on the white dwarf surface. ",
        "introduction": "Cataclysmic variables (CVs) are close binary systems consisting of a late-type star transferring material onto a white dwarf (WD) via Roche-lobe overflow. Magnetic CVs (mCVs) are a small subset of the catalogued CVs ($\\approx 10\\% - 20\\%$, \\citealt{downes}; \\citealt{RKcat}), and fall into two categories: polars (or AM Her types after the prototype system) and intermediate polars (IPs or DQ Her types). The WDs in polars possess such strong magnetic fields that they can synchronise the whole system, yielding $P_{orb}=P_{spin}$. The strong magnetic field in these systems is confirmed by strong optical polarisation and measurments of cyclotron humps (\\citealt{warner}). Accretion in polars is thought to follow the magnetic field lines of the WD straight from the L1 point onto the WD magnetic poles, and no accretion disk is expected (for a review of polars, see \\citealt{cropper}). Another possible class of systems called asynchronous polars (APs) might exist, where the spin and orbital periods are out of synchronisation by only a few percent. It is not known exactly why this is but one suggestion is that these systems are polars which have had a recent nova event, kicking them slightly out of synchronisation (\\citealt{warner}). For IPs, the lack of strong optical polarisation implies a weaker magnetic field, not powerful enough to synchronise the secondary (for a review of IPs, see \\citealt{patterson}). In these systems, material leaving the L1 point usually forms an accretion disc up to the point where the magnetic pressure exceeds the ram pressure of the accreting gas. From this point onwards the accretion dynamics are governed by the magnetic field lines, which channel the material onto the WD magnetic poles. The nature of these systems is confirmed by the detection of coherent X-ray modulations associated with the spin period of the WD.  In the simplest scenario for X-ray production in mCVs, the magnetically channeled accretion column impacts the WD poles producing hard X-rays from thermal bremsstrahlung cooling by free electrons with kT of the order of 10s of keV  (\\citealt{warner, cropper}). The emission is thought to originate in the post-shock region, a region below the shock front created from the impacting accretion column. Softer X-rays are also produced from the absorption and reprocessing of these higher energy photons in the WD photosphere. As a result, both polars and IPs are expected to emit high energy photons, but discrepancies exist between the observed ratio of soft-to-hard X-rays between polars and IPs, with polars showing an excess of observed soft X-rays (\\citealt{lamb}).  \\cite{chanmugan} and others have reported that the total X-ray luminosities of IPs are greater than those of polars by a factor of $\\approx 10$, attributed mainly to the higher accretion rates. Moreover it has been proposed that strong magnetic fields in polars produce a more ``blobby'' flow than in IPs (\\citealt{warner}). These high density ``blobs'' are then able to penetrate within the post-shock region, emitting fewer bremsstrahlung photons and contributing more to the observed X-ray blackbody spectral component, thought to be produced at the base of the post-shock region.  In recent years an increasing number of mCVs have been detected and discovered by hard X-ray telescopes such as {\\it INTEGRAL}/IBIS (\\citealt{barlow,landi}), {\\it Swift}/BAT (\\citealt{brunsch}) and {\\it SUZAKU}/HXD (\\citealt{terada}). Here, we present a newly compiled list of IBIS mCVs, forming part of a larger compilation of hard X-ray selected IPs from all three telescopes mentioned above, allowing us to study the general properties of this hard X-ray selected sample.  This selected sample of CVs will allow us to investigate the observational properties of the $P_{orb}$--$P_{spin}$ plane, and compare them to mCV evolution models. Moreover, we will investigate the spectral properties of our sample, and present for the first time observed spectral hardness correlations observed within IPs as a function of their orbital and spin parameters. ",
        "conclusions": "This paper has presented a catalogue and analysis of a sample of CVs detected in the hard X-ray range ($>$ 17 keV) with IBIS, BAT and HXD. As with previously compiled high-energy samples of CVs, it is shown that most systems are magnetic. Moreover, some of the detected systems are very rare types of objects (2 APs). The sample is dominated by intermediate polars, with only 2 synchronous polars. This suggests the broadband X-ray/gamma-ray spectrum of IPs is harder and more luminous than that of polars. This could be the effect of accretion rate and magnetic field strength, where IPs have higher accretion and weaker magnetic fields relative to polars, depositing a somewhat smoother accretion stream onto the poles. By contrast to the polars, the accretion flow in IPs may therefore not bury itself deep within the post-shock region, producing a harder broadband X-ray/gamma-ray spectrum.  We have shown that only IPs with $P_{spin}/P_{orb}<0.1$ are consistently found by hard X-ray surveys. Moreover, we have examined the observational properties of mCVs in the $P_{orb} - P_{spin}$ plane. We find that the observations are consistent with the theoretical models of \\cite{norton_theo1,norton_theo2} for mCV evolution, where IPs tend to cluster at about $P_{spin}/P_{orb} \\approx 0.1$, and none have yet been observed in the hard X-ray regime above $P_{spin}/P_{orb} \\approx 0.1$. Also observed and predicted is the observation that a very low number of IPs are found in any wavelength range within the synchronicity gap: a region between $P_{spin}/P_{orb} \\approx 0.3$ and $P_{spin}/P_{orb}=1$.  The paper has also presented the first observed correlations between 30-60 keV/17-30 keV hardness and $P_{spin}$, $P_{orb}$ and $P_{spin}/P_{orb}$. The correlations have been statistically tested using Monte Carlo simulations.  In an attempt to explain our result we have suggested that hard X-ray selected IPs are spinning towards their equilibrium, so that their spin period is an indicator of magnetic field strength. This in turn will give rise to a short, but wide, post-shock region for fast spinning WDs (and therefore possessing a relatively weak magnetic field) making their hard X-ray spectra harder. In contrast slowly spinning WDs will have a tall but narrow post-shock region (possessing a relatively high magnetic field), yielding a cooler bremsstrahlung component in the hard X-rays.  All of the observations presented in this paper are consistent with mCV evolution models. It is very likely that hard X-ray missions will continue to increase this sample of mCVs, and it is also expected that more unidentified hard X-ray sources will be identified as IPs with more optical follow-ups. In particular, more observations will allow us to establish if any IPs are detected below the period gap and if any IPs will ever get detected in the IBIS energy range above $P_{spin}/P_{orb} \\approx 0.1$ in order to establish if the Norton accretion models are a plausible explanation to the observed systems. This also implies that hard X-ray selected samples could have their own biases, however more analysis will have consequences on evolution studies of these exotic magnetic systems."
    },
    "0910/0910.0215_arXiv.txt": {
        "abstract": "The IceCube detector is an all-flavor neutrino telescope. For several years IceCube has been detecting muon tracks from charged-current muon neutrino interactions in ice. However, IceCube has yet to observe the electromagnetic or hadronic particle showers or ``cascades'' initiated by charged- or neutral-current neutrino interactions.  The first detection of such an event signature will likely come from the known flux of atmospheric electron and muon neutrinos. A search for atmospheric neutrino-induced cascades was performed using a full year of IceCube data.  Reconstruction and background rejection techniques were developed to reach, for the first time, an expected signal-to-background ratio $\\sim$1 or better. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.1025_arXiv.txt": {
        "abstract": " ",
        "introduction": "Since the discovery of polarized broad emission lines in the archetypal Seyfert 2 galaxy NGC 1068 (Antonucci \\& Miller 1985), the AGN Unified Model (UM) was tested over the entire electromagnetic spectrum. While it has been realized that its original formulation was too simple to explain the large body of observational data, it still provides a useful framework for AGN studies.  Hard X--ray observations of Seyfert galaxies have shown that absorption by circumnuclear gas and dust, presumably with a toroidal geometry, is almost ubiquitous among optically classified Type 2 objects. Deep and wide X--ray surveys, combined with pointed observations of nearby bright AGN, allowed to probe a large interval of luminosities, redshifts and column densities and thus to test the AGN UM in the X--ray band.  A wide range of absorbing column densities is observed among Seyfert 2 and is required to fit the X--ray background ({\\sc xrb}) spectrum (Gilli et al 2007; Treister et al. 2009). Absorption is much less common at high X--ray luminosities with a trend similar to that observed at optical and infrared wavelengths (Simpson 2005, Maiolino et al. 2007). There are also hints of an increasing fraction of obscured AGN towards high redshifts (La Franca et al. 2005; Treister et al. 2006). Besides providing additional tests to the UM, the above trends are most likely closely related to the growth and evolution of Supermassive Black Holes (SMBHs). According to current models (e.g. Hopkins et al. 2008) all the galaxies undergo a phase of heavy, possibly Compton Thick ($N_H > 10^{24}$ cm$^{-2}$) obscuration, strong  accretion and star formation. The census of obscured AGN may thus provide useful insights on our understanding of a key phase of SMBH and galaxy co-evolution. Due to the lack of sensitive X--ray observations above 10 keV, Compton Thick AGN can be efficiently detected and recognized as such only in the local Universe.  In the following we present {\\it Suzaku} observations of a small sample of hard X--ray selected Seyfert 2. The goal is to characterize the physics and geometry of the obscuring gas and, a first step, to investigate their local space density. We also briefly discuss the search for the most obscured sources in various X--ray surveys and the implications for the study of their properties at high redshifts. ",
        "conclusions": ""
    },
    "0910/0910.1213_arXiv.txt": {
        "abstract": "{SMM~J$04542-0301$ is an extended ($\\sim 1 \\arcmin$) submm source located near the core of the cluster MS$0451.6-0305$. It has been suggested that part of its emission arises from the interaction between a LBG and two EROs at $\\rm z\\sim2.9$ that are multiply-imaged in the optical/NIR observations. However, the dramatic resolution difference between the sub-mm map and the optical/NIR images make it difficult to confirm this hypothesis.}{In a previous paper, we reported the detection of 1.4~GHz continuum radio emission coincident with this sub-mm source using VLA archival data. To fully understand the relation between this radio emission, the sub-mm emission, and the optical/IR multiply-imaged sources, we have re-observed the cluster with the VLA at higher resolution.}{The previous archival data has been re-reduced and combined with the new observations to produced a deep ($\\sim$ 10~$\\mu$Jy beam$^{-1}$), high resolution ($\\sim 2\\arcsec$) map centred on the cluster core. The strong lensing effect in the radio data has been quantified by constructing a new lens model of the cluster.}{From the high resolution map we have robustly identified six radio sources located within SMM~J$04542-0301$. The brightest and most extended of these sources (RJ) is located in the middle of the sub-mm emission, and has no obvious counterpart in the optical/NIR. Three other detections (E1, E2 and E3) seem to be associated with the images of one of the EROs (B), although the NIR and radio emission appear to originate at slightly different positions in the source plane. The last two detections (CR1 and CR2), for which no optical/NIR counterpart have been found, seem to constitute two relatively compact emitting regions embedded in a $\\sim 5\\arcsec$ extended radio source located at the position of the sub-mm peak. The presence of this extended component (which contributes 38\\% of the total radio flux in this region) can only be explained if it is being produced by a lensed region of dust obscured star formation in the center of the merger. A comparison between the radio and sub-mm data at the same resolution suggests that E1, E2, E3, CR1 and CR2 are associated with the sub-mm emission.}{The radio observations presented in this paper provide strong observational evidence in favour of the merger hypothesis. However, the question if RJ is also contributing to the observed sub-mm emission remains open. These results illustrate the promising prospects for radio interferometry and strong gravitational lensing to study the internal structure of SMGs.}{} ",
        "introduction": "The detection of the cosmic infrared background (CIB) by the COBE satellite \\citep{PU96.1, HA98.1} established that about half of the total radiation in the universe comes from dust-obscured galaxies that are missing from optical surveys \\citep[see][for a review]{LA05.1}. This population of dusty objects was first resolved by the IRAS and ISO satellites up to $z\\sim1$, and turned out to be dominated by luminous and ultra-luminous IR galaxies (LIRGs and ULIRGs). The step into the high-z universe came with the advent of sub-mm and mm surveys, since the far-infrared (FIR) luminosity peak of high-z obscured galaxies is red-shifted into the sub-mm band \\citep{FR91.1,BL93.1}. As a result, sub-mm galaxies (hereafter SMGs) turned out to be hundreds of times more numerous than galaxies with similar luminosities in the local universe, suggesting that they constitute the dominant contributor of the CIB and cosmic star formation at $z\\gtrsim1$. This illustrates the very important role that SMGs play in the context of galaxy formation and evolution \\citep[see][for a review on (U)LIRGs and SMGs]{LO06.1}.  One decade after their discovery with SCUBA\\footnote{Submillimeter Common-User Bolometer Array, decommissioned in September 2005 from the James Clerk Maxwell Telescope on Mauna Kea} \\citep{SM97.1,HU98.1,BA98.1,EA99.1}, it is generaly acepted that SMGs are heavily dust-obscured galaxies at high redshift ($2<z <3$) with ULIRG-like luminosities ($L_{\\rm FIR} \\sim 10^{12}~ L_\\odot$) and star formation rates of the order of 1000 $M_\\odot~yr^{-1}$. This enormous bolometric luminosity seems to be dominated by star formation processes induced by galaxy interactions/mergers, although a good fraction ($\\sim 30-50 \\%$) of SMGs also host AGN activity \\citep{AL05.1}. First estimates of their physical properties indicate that SMGs are massive, gas rich systems ($M_{\\rm gas}\\sim10^{10}-10^{11}~M_\\odot$) in which the starburst region has a typical scale in the range $1-8$~kpc \\citep{NE03.1, CH04.1, GR05.1, TA06.1, WA07.1, BI08.1}. The available evidence also suggests that SMGs might be the progenitors of massive local ellipticals \\citep{LI99.1, SM02.1, SM04.1, WE03.1, GE03.1, AL03.1,AL05.1, SW06.1}.  However, these general properties of SMGs are based on the study of the very brightest examples of this class of object ($\\rm S_{850\\mu m}\\gtrsim2$~mJy), and may not be representative of the entire SMG population. In fact, according to \\cite{KN08.1}, the dominant contribution to the sub-mm extragalactic background comes from the fainter (sub-)mJy sources that cannot usually be detected due to the confusion noise of current instruments ($\\rm S_{850\\mu m}\\sim 2$~mJy). The only way in which it has been possible to push bellow this sensitivity limit, is by using the lensing magnification provided by massive clusters of galaxies to increase the effective resolution of SCUBA. This approach has improved the sensitivity of sub-mm maps by factors of a few with respect to blank field surveys, although just a handfull of faint SMGs have being identified so far \\citep{SM02.1, CO02.1, KN08.1}. Since so little is known about these intrinsically faint sources, it is crucial to further investigate their properties (e.g. spectral energy distributions, morphologies, redshifts, etc.) and assess whether they are different from the observed properties of brighter SMGs.  A promissing strategy to gather information about faint SMGs is to study members that are multiply imaged by clusters of galaxies. In this cases, the magnification factor can go up to 30 (or more), providing not only the opportunity to detect but also to spatially resolve the morphologies and internal dynamics of faint SMGs at a level of detail far greater than would otherwise be possible \\citep[see][for an example of this technique in the optical]{SW07.1}. To date, only one multiply-imaged faint SMGs has been confirmed: SMM~J$16359+6612$, located near the core of the cluster A2218 \\citep{KN04.1,KN05.1,GA05.1,KN08.1,KN09.1}. There are, however, two other clusters which seem to host multiply-imaged SMGs: A1689 \\citep{KN08.1} and MS$0451.6-0305$ \\citep{CH02.1, BO04.1,BE07.1}.   \\begin{figure} \\centering \\includegraphics[width=9cm,angle=0]{12903fg1.ps} \\caption{\\textbf{Summary of the results presented in \\cite{BO04.1}.} SCUBA $850~\\mu$m contour map of SMM~J$04542-0301$ superimposed upon a HST image of the center of the cluster MS$0451.6-0305$. Highlighted with rectangles are an optical arc (ARC1) and its counter image (ARC1 ci), produced by a LBG at $z_{\\rm spect}= 2.911$. The red circles indicate the positions of five NIR sources that have been interpreted as multiple images produced by two EROs (ERO B, lensed as B1/B2/B3, and ERO C, lensed as C1/C2/C3) at $z_{\\rm model}=2.85 \\pm 0.1$. Assuming $z=2.9$ for both the LBG and the EROs, their predicted positions in the source plane would be located withing a region of $\\rm ~\\sim 10~kpc$, suggesting that they constitute a merger.} \\label{intro} \\end{figure}   The case of MS$0451.6-0305$ is particularly interesting, since the ``sub-mm source'' (SMM~J$04542-0301$) is an elongated ($\\sim 1 \\arcmin$) region of $850~\\mu$m emission which is coincident with an optical arc and five NIR sources (see Fig. \\ref{intro}).  While the optical arc is the result of a strongly-lensed Lyman Break Galaxy (LBG) at $z_{\\rm spect}=2.911$, a lens model of the cluster predicts that the set of NIR sources could be produced by two triply-imaged EROs\\footnote{Extremely Red Objects, photometrically defined as $R-K > 5.3$} located at almost the same redshift ($z_{\\rm model}=2.85 \\pm 0.1$). The model also indicates that the ERO pair and the LBG may constitute a merger at $z=2.9$ in the source plane, with their interaction likely being at the origin of the observed sub-mm emission \\citep[][B04 hereafter]{BO04.1}. Unfortunately, the low resolution ($\\sim 15\\arcsec$ at $850~\\mu$m) and poor positional accuracy ($\\sim 2\\arcsec-3\\arcsec$ rms) of SCUBA, makes it very difficult to confirm the link between SMM~J$04542-0301$ and the proposed optical/NIR sources. Moreover, the emission coming from the north-eastern and central regions of the sub-mm emission cannot be reproduced by the arc and the NIR sources (see Fig. 7 in B04), suggesting that it might arise via other sources.   A possible way to overcome this resolution problem is by taking advantage of the observed correlation between the radio synchrotron and FIR emission in star-forming galaxies \\citep{VK73.1,CO82.1,HE85.1,GA02.1,AP04.1,BE06.1}, which also seems to hold for SMGs out to $z\\sim3$ \\citep{KO06.1,VL07.1,IB08.1,MI09.1}. Thanks to this FIR-radio correlation, radio interferometric observations can be used as a high-resolution proxy for the rest-frame FIR emission observed in the sub-mm. This approach has been extensively used to pinpoint the position of SMGs in order to identify their faint optical counterparts \\citep[e.g.][]{IV00.1,BA00.1,CH05.1,IV07.1}.    In \\cite{BE07.1}, we reported the detection of 1.4~GHz radio emission coincident with SMM~J$04542-0301$ using VLA\\footnote{Very Large Array interferometer of the National Radio Astronomy Observatory (NRAO)} archival data. Part of this radio emission is located in the region between the optical arc and the ERO images, which is consistent with the interacting region of the hypothetical merger being the source of the observed radio and sub-mm emission. We also detected bright radio emission in the central region of SMM~J$04542-0301$, although it is not clear if this emission is produced by a high-z lensed object or AGN activity associated with a cluster member.  In this paper, we present a higher resolution 1.4~GHz radio map of the center of the cluster MS$0451.6-0305$ (MS0451 hereafter), obtained after combining the previous VLA archival data (BnA configuration) with new VLA high resolution observations (A configuration).  The details of both sets of observations are presented in Sect. 2, together with a description of the data analysis procedure, which has been considerably improved with respect to \\cite{BE07.1}. Section 3 is dedicated to the analysis of the final combined data set, from which we characterise the compact and extended radio emission detected within SMM~J$04542-0301$. Identification of optical (HST/ACS F814W) and NIR (SUBARU/CISCO K' band) counterparts for the radio detections, as well as their connection with the sub-mm emission, is discussed in Sect. 4. In Sect. 5 we describe a new lens model of the cluster MS$0451.6-0305$, built to investigate the lensed nature of the compact radio detections. A discussion about the merger scenario proposed by \\cite{BO04.1} including the results from the radio observations are presented in Sect. 6. Summary and conclusions are presented in Sect. 7.  The adopted cosmology corresponds to a $\\Lambda$CDM model with $\\Omega_{\\rm m}=0.28$, $\\Omega_{\\lambda}=0.72$ and $h_{0}=73$ \\citep{SP07.1}. ",
        "conclusions": ""
    },
    "0910/0910.4600_arXiv.txt": {
        "abstract": "We report the discovery of four very bright, strongly-lensed galaxies found via systematic searches for arcs in Sloan Digital Sky Survey Data Release 5 and 6. These were followed-up with spectroscopy and imaging data from the Astrophysical Research Consortium 3.5m telescope at  Apache Point Observatory and found to have redshift $z>2.0$. With isophotal magnitudes $r = 19.2 - 20.4$ and $3\\arcsec$-diameter magnitudes $r = 20.0 - 20.6$, these systems are some of the brightest and highest surface brightness lensed galaxies known in this redshift range. In addition to the magnitudes and redshifts, we present estimates of the Einstein radii, which range from $5.0 \\arcsec$ to $12.7 \\arcsec$, and use those to derive the enclosed masses of the lensing galaxies. ",
        "introduction": "Strong gravitational lenses provide an opportunity for studying properties of distant galaxies. Because surface brightness is unchanged during lensing, magnification of the image provides amplification of the image flux and allows studies of details that would otherwise be too faint for detailed ground-based investigation. Models of the lens systems also provide information on the mass distribution, including dark matter, of the lenses themselves.  Recent wide area optical surveys have enabled identification of new samples of strong lens systems. The Sloan Digital Sky Survey (SDSS)~\\cite{york} spectroscopic data has been used~\\cite{slacsfull,ols-lens} to identify $\\sim 140$ strong lensing candidates in systems where multiple spectra appear in a single $3\\arcsec$-wide spectral fiber. The SDSS imaging data provides the opportunity to identify a complementary set of strong lens systems that have a larger lens-arc separation $(> 3.0\\arcsec)$. An imaging search of 240 rich clusters~\\cite{Hennawi} yielded 16 strong lens systems with large ($> 10\\arcsec$) radius, where the lensing interpretation is based on morphology, and 21 additional systems for which the lensing interpretation is less certain.  Several more large arcs in the SDSS data have been described in recent papers~\\cite{belokurov,wen}. The CFHTLS field  and COSMOS field have yielded 40~\\cite{Cabanac} and 70~\\cite{faure,jackson} strong lens candidate systems.  Relatively few arc systems have been reported where the lensed galaxy is known to have redshift $z > 2$.  Some systems with high surface brightness have been reported. The brightest Lyman Break Galaxy (LBG) currently known is SDSS J002240.91+143110.4 ``The 8 o'clock arc''~\\cite{8oclock} with a  redshift $z=2.73$ and magnitude $r=19.2$. Other bright systems include SDSS J120602.09+514229.5 ``The Clone''~\\cite{clone} with redshift $z=2.0$ and magnitude $r=19.8$, MS 1512-cB58~\\cite{cB58} with redshift $z=2.73$ and magnitude $r=20.4$, and the system~\\cite{z307} lensed by MACS J2135.2-0102 with redshift $z=3.07$.   Motivated by our discovery of ``The 8 o'clock arc'', we initiated a program, the Sloan Bright Arcs Survey, to identify and confirm lensed high-redshift galaxies. We recently reported~\\cite{kubo} the discovery of six spectroscopically confirmed systems with redshifts $z = 0.4$ to $1.4$. In this paper we report the discovery and confirmation of four more systems with redshift $z>2$.  This paper is organized as follows. In $\\S 2$ we describe the strong-lens arc search samples and candidate identification procedure. In $\\S 3$ we describe our follow-up imaging and spectroscopy procedure. In $\\S 4$ we describe the results from the follow-up and use those results to calculate properties of the lensed galaxies. Finally, in $\\S 5$, we describe further follow-up measurements underway and summarize. ",
        "conclusions": "Our strong lens search has yielded four systems with lensed galaxies that have redshift $z>2.0$. The lens redshifts are presented, based on the SDSS spectroscopic data. We measured the source galaxy redshifts using the ARC 3.5m DIS spectrograph. Further imaging using the ARC 3.5m imager SPICAM revealed more detail about the lensed systems.  The arcs from the lensed galaxies have Einstein radius between $5.0\\arcsec$ and $12.7\\arcsec$ and are quite bright. A simple SIS model was used to provide estimates of the mass enclosed within the Einstein radius of the lensing galaxies.  From the data in Table~\\ref{tab-sys}, we see that these systems have approximate total isophotal magnitudes $r = 19.2 - 20.4$. For comparison with other lensed systems, we use the $3\\arcsec$-diameter magnitudes and those are $r = 20.0 - 20.6$. Thus these systems are about one magnitude fainter than the brightest lensed LBG known, the 8 o'clock arc ($r = 19.2$)~\\cite{8oclock}, but are quite similar to the few other very bright $z > 2$ lensed systems, such as MS 1512-cB58 ($r = 20.4$)~\\cite{cB58} and the Clone ($r = 19.8$) ~\\cite{clone}.  Likewise, the brightest knots in these systems also have very high surface brightness, $\\sim 23.5 \\; {\\rm mag \\; arcsec}^{-2}$ in $r$, similar to the Clone, and about one magnitude fainter than the 8 o'clock arc. The systems in our sample are thus some of the brightest lensed galaxies known at $z > 2$ and will be amenable to detailed studies of their individual properties.  Combined with our previous results~\\cite{8oclock, clone, kubo} the Sloan Bright Arcs Survey has now spectroscopically confirmed  12 strong lensing systems in SDSS DR5 and DR6. There are follow-up observations of all of these systems scheduled for either Cycle 16 Supplemental data or for Cycle 17 of the Hubble Space Telescope. These observations will provide high-quality images useful for detailed lens and source modeling. Additional follow-up observations obtained using WIYN 3.5m telescope at Kitt Peak National Observatory may provide wide-field images of the lensing galaxy that can be used to characterize the galaxy cluster environment.  We continue to follow up additional promising lens candidates."
    },
    "0910/0910.3166_arXiv.txt": {
        "abstract": "\\emph{Swift} captured for the first time a smoothly rising X-ray re-brightening of clear non-flaring origin after the steep decay in a long gamma-ray burst (GRB): GRB\\,081028. A rising phase is likely present in all GRBs but is usually hidden by the prompt tail emission and constitutes the first manifestation of what is later to give rise to the shallow decay phase. Contemporaneous optical observations reveal a rapid evolution of the injection frequency of a fast cooling synchrotron spectrum through the optical band, which disfavours the afterglow onset (start of the forward shock emission along our line of sight when the outflow is decelerated) as the origin of the observed re-brightening. We investigate alternative scenarios and find that the observations are consistent with the predictions for a narrow jet viewed off-axis. The high on-axis energy budget implied by this interpretation suggests different physical origins of the prompt and (late) afterglow emission. Strong spectral softening takes place from the prompt to the steep decay phase: we track the evolution of the spectral peak energy from the $\\gamma$-rays to the X-rays and highlight the problems of the high latitude and adiabatic cooling interpretations. Notably, a softening of both the high and low spectral slopes with time is also observed. We discuss the low on-axis radiative efficiency of GRB\\,081028 comparing its properties against a sample of \\emph{Swift} long GRBs with secure $E_{\\rm{\\gamma,iso}}$ measurements. ",
        "introduction": "Gamma ray bursts (GRBs) are transient events able to outshine the $\\gamma$-ray sky for a few seconds to a few minutes. The discovery of their optical \\citep{Paradijs97} and X-ray \\citep{Costa97} long-lasting counterparts represented a breakthrough for GRB science. Unfortunately, due to technological limitations, the X-ray observations were able to track the afterglow evolution starting hours after the trigger: only after the launch of the \\emph{Swift} satellite in 2004 \\citep{Gehrels04} was this gap between the end of the prompt emission and several hours after the onset of the explosion filled with X-ray observations. A canonical picture was then established (see e.g., \\citealt{Nousek06}), with four different stages describing the overall structure of the X-ray afterglows: an initial steep decay, a shallow-decay phase, a normal decay and a jet-like decay stage. Erratic flares are found to be superimposed mainly to the first and second stage of emission. An interesting possibility is that the four light-curve phases instead belong to only two different components of emission (see e.g., \\citealt{Willingale07}): the first, connected to the activity of the central engine giving rise to the prompt emission, comprises the flares (\\citealt{Chincarini07} and references therein) and the steep-decay phase; the second is instead related to the interaction of the outflow with the external medium and manifests itself in the X-ray regime through the shallow, normal and jet-like decay. Observations able to further characterise the two components are therefore of particular interest.  The smooth connection of the X-ray steep decay light-curve phase with the prompt $\\gamma$-ray emission strongly suggests a common physical origin (\\citealt{Tagliaferri05}; \\citealt{Obrien06}): the high latitude emission (HLE) model (\\citealt{Fenimore96}; \\citealt{Kumar00}) predicts that steep decay photons originate from the delay in the arrival time of prompt emission photons due to the longer path length from larger angles relative to our line of sight, giving rise to the $\\alpha=\\beta+2$ relation (where $\\alpha$ is the light-curve decay index and $\\beta$ is the spectral energy index). No spectral evolution is expected in the simplest formulation of the HLE effect in the case of a simple power-law prompt spectrum. Observations say the opposite: significant variations of the photon index have been found in the majority of GRBs during the steep decay phase (see e.g., \\citealt{Zhang07}); more than this, the absorbed simple power-law (SPL) has proved to be a poor description of the spectral energy distribution of the steep decay phase for GRBs with the best statistics\\footnote {The limited $0.3 -10 \\,\\rm{keV}$ spectral coverage of the Swift X-Ray Telescope, XRT \\citep{Burrows05}, and the degeneracy between  the variables of the spectral fit can in principle lead to the identification of an SPL behaviour in intrinsically non-SPL spectra with poor statistics.}. A careful analysis of these events has shown their spectra to be best fit by an evolving Band function \\citep{Band93}, establishing the link between steep decay and prompt emission photons also from the spectral point of view (see e.g., GRB\\,060614, \\citealt{Mangano07}; GRB\\,070616, \\citealt{Starling08}): caused by the shift of the Band spectrum, a temporal steep decay phase and a spectral softening appear simultaneously (see e.g. \\citealt{Zhang09}, \\citealt{Qin09}). In particular, the peak energy of the $\\nu F_{\\nu}$ spectrum is found to evolve to lower values, from the $\\gamma$-ray to the soft X-ray energy range. Both the low (as observed for GRB\\,070616) and the high-energy portion of the spectrum are likely to soften with time, but no observation is reported to confirm the high energy index behaviour during the prompt and steep decay phase. The observed spectral evolution with time is an invaluable footprint of the physical mechanisms at work: observations able to constrain the behaviour of the spectral parameters with time are therefore of primary importance.  By contrast, no spectral evolution is observed in the X-ray during the shallow decay phase (see e.g. \\citealt{Liang07}) experienced by most GRBs between $\\sim 10^2$ s and $10^3 -10^ 4$ s. An unexpected discovery of the \\emph{Swift} mission, the shallow decay is the first light-curve phase linked to the second emission component. A variety of theoretical explanations have been put forward. The proposed models include: energy injection (\\citealt{Panaitescu06}; \\citealt{Rees98}; \\citealt{Granot06}; \\citealt{Zhang06}); reverse shock (see e.g., \\citealt{Genet07}); time dependent micro-physical parameters (see e.g. \\citealt{Granot06b}; \\citealt{Ioka06}); off-axis emission \\citep{Eichler06}; dust scattering \\citep{Shao07}. The predictions of all these models can only be compared to observations tracking the flat and decay phase of the second emission component, since its  rise is usually missed in the X-ray regime, being hidden by the tail of the prompt emission.  GRB\\,081028 is the first and unique event for which \\emph{Swift} was able to capture the rise of the second emission component\\footnote {There are a handful of long GRBs detected by \\emph{Swift} with a possible X-ray rise of non-flaring origin. Among them: GRB\\,070328, \\cite{Markwardt07}; GRB\\,080229A, \\cite{Cannizzo08}; GRB\\,080307, \\cite{Page09} (see \\citealt{Page09} and references therein). However, in none of these cases has an X-ray steep decay been observed. A smooth rise in the X-rays has been observed in the short GRB\\,050724.}: the time properties of its rising phase can be constrained for the first time while contemporaneous optical observations allow us to track the evolution of a break energy of the spectrum through the optical band. GRB\\,081028 is also one of the lucky cases showing a spectrally evolving prompt emission where the evolution of the spectral parameters can be studied from $\\gamma$-rays to X-rays, from the trigger time to $\\sim 1000$ s. A hard to soft spectral evolution is clearly taking place beginning with the prompt emission and extending to the steep decay phase, as already found for other \\emph{Swift} GRBs (GRB\\,060614, \\citealt{Mangano07}; GRB\\,070616, \\citealt{Starling08}, are showcases in this respect). Notably, for GRB\\,081028 a softening  of the slope of a Band function  \\citep{Band93} above $E_{\\rm p}$ is also observed.  The paper is organised as follows: \\emph{Swift} and ground-based observations are described in Sect. \\ref{sec:observations}; data reduction and preliminary analysis are reported in Sect. \\ref{Sec:datared}, while in Sect. \\ref{sec:analysis} the results of a detailed spectral and temporal multi-wavelength analysis are outlined and discussed in Sect. \\ref{sec:discussion}. Conclusions are drawn in Sect. \\ref{sec:conclusion}.  The phenomenology of the burst is presented in the observer frame unless otherwise stated. The convention $F_{\\nu}(\\nu,t)\\propto \\nu^{-\\beta}t^{-\\alpha}$ is followed, where $\\beta$ is the spectral energy index, related to the spectral photon index $\\Gamma$ by $\\Gamma=\\beta+1$. All the quoted uncertainties are given at 68\\% confidence level (c.l.): a warning is added if it is not the case. The convention $Q_{x}=Q/10^{x}$ has been adopted in cgs units unless otherwise stated. Standard cosmological quantities have been adopted: $H_{0}=70\\,\\rm{km\\,s^{-1}\\,Mpc^{-1}}$, $\\Omega_{\\Lambda }=0.7$, $\\Omega_{\\rm{M}}=0.3$.\\\\ ",
        "conclusions": "\\label{sec:conclusion}  The 0.3-10 keV X-ray emission of GRB\\,081028 consists of a flat phase up to $\\sim300$ s (the XRT is likely to have captured the prompt emission in the X-ray energy band) followed by a steep decay with flares superimposed extending to $\\sim7000$ s (component 1). The light-curve then shows a re-brightening which starts to rise at $t\\sim 8000\\,\\rm s$ and peaks around $20\\,\\rm ks$ (component 2). The different spectral and temporal properties strongly characterise the XRT signal as due to two distinct \\emph{emission} components. However, their further characterisation as emission coming from \\emph{physically} distinct \\emph{regions} is model dependent.   The strong hard-to-soft evolution characterising the prompt and steep decay phase of GRB\\,081028 from trigger time to 1000 s is well modelled by a shifting Band function: the spectral peak energy evolves to lower values, decaying as $E_{\\rm peak}\\propto t^{-7.1\\pm0.7}$ or $E_{\\rm peak}\\propto (t-t_0)^{-4.2\\pm2.4}$ when the zero-time of the power-law is allowed to vary: the best fit constrains this parameter to be $t_0=109\\pm89$ s. In either case our results are not consistent with the $\\propto t^{-1}$ behaviour predicted by the HLE in its simplest formulation. While a more realistic version of this model might still account for the observed $E_{\\rm peak}$ evolution, other possibilities must be investigated as well: the adiabatic expansion cooling of the $\\gamma$-ray source predicts a steeper than observed light-curve decay and is therefore unlikely. While the peak is moving, a softening of both the low and high-energy portions of the spectrum is clearly detected. The failure of both the curvature effect and the adiabatic cooling argues against the abrupt switch-off of the GRB source after the prompt emission and suggest the continuation of the central engine activity during the steep decay. An off-axis explanation may reconcile the high latitude emission or the adiabatic expansion cooling models with the data. This will be explored in a future work.  GRB\\,081028 has afforded us the unprecedented opportunity to track a smoothly rising X-ray afterglow after the steep decay: the rising phase of the emission component later accounting for the shallow light-curve phase is usually missed, being hidden by the steep decay which is the tail of the prompt emission both from the spectral and from the temporal point of view. The peculiarity of GRB\\,081028 lies in a small overlap in time between the steep decay and the following re-brightening caused by an unusual delay of the onset of the second component of emission. Contemporaneous optical data allow the evolution of the SED during the re-brightening to be constrained: the spectral distribution is found to be best described by a photo-electrically absorbed smoothly broken power-law with a break frequency evolving from $1.6\\times 10^{15}\\,\\rm Hz$ downward to the optical band. The break frequency can be identified with the injection frequency of a synchrotron spectrum in the fast cooling regime evolving as $\\nu_{\\rm b}\\propto t^{-2.6\\pm0.2}$. The intrinsic optical absorption is found to satisfy $A_{V,z}<0.22$.  The observed break frequency scaling is inconsistent with the standard predictions of the onset of the forward shock emission even if this model is able to account for the temporal properties of the X-ray re-brightening (note that in this context the delay of the second emission component is due to a lower than usual fireball Lorentz factor or external medium density). Alternative scenarios have therefore been considered. While a dust scattering origin of the X-ray emission is ruled out since we lack observational evidence for a non-negligible dust extinction and strong spectral softening, a reverse shock origin cannot be excluded. However, this can be accomplished only by requiring non-standard burst parameters: the ejecta should have a tail of Lorentz factors decreasing to low values; $\\epsilon_{e}$ should be near equipartition; only a small fraction $\\xi_{e}\\sim 10^{-2}$ of electrons should contribute to the emission.  The predictions of the off-axis model have been discussed in detail: according to this model a peak of emission is expected when the beaming cone widens enough to engulf the line of sight. The delayed onset of the second emission component is not a consequence of unusual intrinsic properties of the GRB outflow but is instead an observational artifact, due to the off-axis condition. The observed evolution of $\\nu_{\\rm b}$ is consistent with the expected evolution of the injection frequency of a fast cooling synchrotron spectrum for $0\\lesssim k\\lesssim2$. We interpret the light-curve properties as arising from an off-axis view, with $\\theta\\sim3\\Delta \\theta$ and $\\theta\\sim0.03(\\frac{E_{54}}{n_{0}})^{-1/8}$ for $k$=0 (or $\\theta\\sim0.03(\\frac{E_{54}}{A_{*}})^{-1/4}$ for $k$=2), $\\theta$ being the angle from the outer edge of the jet and $\\Delta\\theta$ the jet opening angle. %  In this scenario, the peculiarity of GRB\\,081028, or the reason why we do not observe more GRB\\,081028-like events, may be attributed to the following reasons. Since GRB\\,081028 is a particularly bright (and therefore rare) event when viewed on-axis (with high on-axis $E_{\\rm{iso}}$ and $L_{\\rm{iso}}$ values), it is detectable by an off-axis observer even at the cosmological distance implied by its redshift z = 3.038. In addition, GRB\\,081028 appears to be characterized by a particularly narrow jet, for which the ratio of the detectable off-axis solid angle to on-axis solid angle is larger than for wider (but otherwise similar) jets. Finally, GRB\\,081028 might have a peculiar angular structure that is not representative of most GRBs, which would undermine the drawing of statistical conclusions under the assumption of a similar angular structure for most or all GRB jets.    The radiative efficiency is one of the key parameters in GRB science: a precise estimate of this parameter would allow one to distinguish between different models put forward to explain the observed emission. For the on-axis model, with $\\epsilon_{\\gamma}\\sim 10^{-2}$, the GRB\\,081028 efficiency turns out to be lower than the values obtained by FP06 and LZ04 for a sample of pre-\\emph{Swift} GRBs: this directly implies that instead of having released as much energy in the prompt emission as most bursts of the two samples, GRB\\,081028 has a much greater kinetic energy injected in the outflow. Figure \\ref{Fig:EisoLisoRest} clearly shows that this conclusion is likely to be extended to other \\emph{Swift} bursts with secure $E_{\\gamma,\\rm{iso}}$ measurement. This picture changes if we consider the off-axis interpretation: if the deceleration time is much longer than the prompt duration  the prompt and afterglow emission are consistent with originating from the same physical component and the efficiency of the burst is comparable to most bursts; if instead the deceleration time is close to the end of the prompt emission, then the on-axis isotropic energy output would imply an extremely high efficiency of $99\\%$ which is very hard to explain. This suggests that the prompt and afterglow emission come from different physical components.   GRB\\,081028 demonstrates the evolution of GRB spectral properties from the onset of the explosion to $\\sim10^{6}\\,\\rm s$ after trigger and shows that this is likely to be attributed to two distinctly contributing components of emission. These can be constrained only by prompt, broad-band coverage and good time resolution observations."
    },
    "0910/0910.1480_arXiv.txt": {
        "abstract": "We present a set of predictions for weak lensing correlation functions in the context of modified gravity models, including a prescription for the impact of the nonlinear power spectrum regime in these models. We consider the DGP and $f(R)$ models, together with dark energy models with the same expansion history. We use the requirement that gravity is close to GR on small scales to estimate the non-linear power for these models. We then calculate weak lensing statistics, showing their behaviour as a function of scale and redshift, and present predictions for measurement accuracy with future lensing surveys, taking into account cosmic variance and galaxy shape noise. We demonstrate the improved discriminatory power of weak lensing for testing modified gravities once the nonlinear power spectrum contribution has been included. We also examine the ability of future lensing surveys to constrain a parameterisation of the non-linear power spectrum, including sensitivity to the growth factor $\\gamma$. ",
        "introduction": "Consistent observational evidence from various cosmological probes shows that the Universe is currently undergoing a period of accelerated expansion. The observed expansion history can be explained using some form of dark energy or a cosmological constant; however this cosmological constant cannot be explained with current particle physics due to its very small value. An alternative approach is to invoke a modification of gravity; there are many different ways that gravity and/or the equation of state of the dark energy can be modified to allow for the expansion history observed. This makes it impossible to differentiate between the effects of modified gravity and dark energy by measuring the background expansion history alone. However, modifying gravity also produces a distinct growth rate of structure; thus the expansion history and growth history together can be used to distinguish between various models of gravity. This consistency relation to test GR has been proposed and explored by many papers \\citep{Uzan:2000mz,Lue:2003ky,Lue:2004rj,Ishak:2005zs,Kunz:2006ca,Chiba:2007rb,Wang:2007fsa,Bertschinger:2008zb,Jain:2007yk,Daniel:2008et,Song:2008vm}. Upcoming weak lensing surveys such as DES\\footnote[1]{http://www.darkenergysurvey.org}, Pan-STARRS\\footnote[2]{http://pan-starrs.ifa.hawaii.edu} and LSST\\footnote[3]{http://www.lsst.org}, and future space surveys such as Euclid\\footnote[4]{http://www.ias.u-psud.fr/imEuclid}, will allow a combination of growth of structure and expansion history to be probed to considerably higher precision, which will allow many gravity models to be excluded.  There has been a great deal of work showing how to use weak lensing to discriminate between different gravity models; however this has been restricted to probing the linear regime of the matter power spectrum \\citep{Afshordi:2008rd,Schmidt:2008hc,Song:2008xd,Thomas:2008tp,Tsujikawa:2008in,Zhao:2009fn,Zhao:2008bn} or uses methods that do not obtain GR at small scales \\citep{Knox:2006fh,Heavens:2007ka,Yamamoto:2007gd,Amendola:2007rr}. The non-linear regime provides much of the power for lensing and can be most easily probed by current and upcoming lensing surveys. This paper examines the effect of including the non-linear regime in modified gravity lensing predictions, including the small-scale GR limit, to see how useful weak lensing will be overall when trying to determine the correct model of gravity. First we look at DGP and $f(R)$ gravity models as examples, and investigate weak lensing's ability to differentiate between these models and dark energy models. We then take a more phenomenological point of view, by parameterising the shape of the matter power spectrum and examining the sensitivity of weak lensing observables to changes to the matter distribution when the expansion history is the same for each model considered. Using these parameters we show how strongly a ground-based survey similar to DES and a space-based survey such as Euclid will be able to discriminate between different growth histories with identical expansion histories.  This paper is organised as follows. In section \\ref{lensmodgrav} we briefly describe the DGP and $f(R)$ models of gravity and how they compare with dark energy models. We describe how we calculate matter power spectra for these models, including the GR small-scale limit. We also describe how we proceed to calculate weak lensing observables from these power spectra. In section \\ref{results} we present the resulting lensing correlation functions, including realistic errors for future surveys taking into account shape measurement noise and cosmic covariance. In section \\ref{parameterisation} we take the alternative approach of parameterising the non-linear power spectrum, and we investigate how sensitive weak lensing is to these parameters which go beyond the usual growth parameter. We present our conclusions in section \\ref{conclusions}.  Throughout this paper we will use a flat cosmology with the WMAP5+SNe+BAO best fit cosmological parameters, which are determined by the background evolution of the Universe. We use a $\\Lambda$CDM background for both $\\Lambda$CDM and $f(R)$, in which case we take $n_{\\rm s}=0.96$, $h=0.71$, $\\Omega _{\\rm m}=0.27 \\pm 0.02$ and $\\sigma_8=0.81 \\pm 0.03$ \\citep{Komatsu:2008hk}. When we use a DGP background, we have $n_{\\rm s}=0.998$, $h=0.66$, $\\Omega _{\\rm m}=0.26 \\pm 0.02$ \\citep{Fang:2008kc} giving a $\\sigma_8=0.66 \\pm 0.03$ for an equivalent $\\Lambda$CDM model. ",
        "conclusions": "In this paper we have presented weak lensing predictions for modified gravity models, including the non-linear regime of the power spectrum.  We have shown how the power spectrum is calculated for DGP, $f(R)$ and QCDM models, using the fitting function of \\cite{Hu:2007pj} to explore deep into the non-linear regime, while including the fact that gravities should tend towards GR on small scales.  We have calculated the total shear power spectrum given the modified gravity power spectrum, and have shown that this will be measured with high signal-to-noise with future lensing surveys such as Euclid and DES. We have taken into account the cosmic covariance in addition to the noise due to the intrinsic shapes of galaxies.  We have shown that there is substantial additional discriminatory power between modified gravity models which is now afforded to us by the inclusion of the nonlinear power regime. We have also shown that using only the \\cite{Smith:2002dz} formula without any attempt to obtain the GR non-linear power spectrum on small scales leads to an overestimatation in the ability of future surveys to differentiate between different growth histories.  We have parameterised the dark matter power spectrum using the growth factor $\\gamma$ and the parameters in the non-linear fitting function to see how well a ground-based survey similar to DES, and a space-based survey such as Euclid, will be able to put constraints on these. We have compared the results from this parameterisation with results obtained from using only linear scales and have shown the constraint on $\\gamma$ to be much tighter in the former case."
    },
    "0910/0910.1355_arXiv.txt": {
        "abstract": "We present a follow-up study of a series of papers concerning the role of close interactions as a possible triggering mechanism of the activity of AGN and starburst (SB) galaxies. We have already studied the close ($\\leq 100 \\; h^{-1}$kpc) and the large scale ($\\leq$ 1 $h^{-1}$ Mpc) environment of Sy1, Sy2 and Bright IRAS galaxies and their respective control samples (Koulouridis et al.). The results led us to the conclusion that a close encounter appears capable of activating a sequence where a normal galaxy becomes first a starburst, then a Sy2 and finally a Sy1 galaxy. However since both galaxies of an interacting pair should be affected, we present here optical spectroscopy and X-ray imaging of the neighbouring galaxies around our Seyfert and BIRG galaxy samples. We find that more than 70\\% of all neighbouring galaxies exhibit thermal or/and nuclear activity (namely enhanced star formation, starbursting and/or AGN) and furthermore we discovered various trends regarding the type and strength of the neighbour's activity with respect to the activity of the central galaxy, the most important of which is that the neighbours of Sy2 are systematically more ionized, and their straburst is younger, than the neighbours of Sy1s. Our results not only strengthen the link between close galaxy interactions and activity but also provide more clues regarding the evolutionary sequence inferred by previous results. ",
        "introduction": "Since the discovery of Active Galactic Nuclei (AGN) significant effort has been put in the attempt to reveal their nature, but we still lack a full understanding of all the aspects of activity. We are almost certain of the existence of a massive black hole (MBH) in the centre of active and even non active (Ferrarese and Merrit 2000; Gebhardt 2000) galaxies (including our own; e.g. Bozza \\& Mancini 2009). In active galaxies accretion of material into this BH is responsible for the excess nuclear emission that we see in the galaxy spectrum. On the other hand, the triggering mechanism and the feeding of the black hole, the structure and the physics of the accretion disk and of the torus, which is predicted by the unified scheme (Antonucci et al. 1993), the connection with star formation and the role of the AGN feedback, the origin of the jets in radio loud AGN galaxies and their influence on the IGM and ICM and even the mechanism that produces the observed IR, X-ray, and gamma-ray emission, are only partially understood. The unification model itself has not been able to fully explain all the AGN phenomenology (in particular, the role of interactions on induced activity; Koulouridis et al. 2006a,b and references therein).  The lack of detailed knowledge of key aspects of the AGN mechanism leaves us with many scattered pieces of information. Theory is unable in most cases to explain observations, and observations cannot resolve the galactic nuclei to confirm theories. Radio-loud and radio-quiet AGN, QSO type I and II, Sy1 and Sy2 galaxies, LINERs, transition galaxies between different states (TOs) and starforming galaxies, are some of the pieces we are called to unify. Examination of the properties of the host galaxies of the different types of AGN and their environment up to several hundred kpc, can give us valuable information about the central engine. In addition, the availability of large, automatically constructed, galaxy catalogues nowadays, like the SDSS, can provide the necessary statistical significance for these type of analyses. Usually the bigger the sample the less the control we have over the spectral details of the individual galaxies and therefore the less control over the interpretation of the output results. It is dangerous to attempt to draw conclusions from analyses of large samples, that include different kind of objects, without at least having hints regarding the answers to basic questions: for example, is the Unification scheme valid for all cases of type 1 and type 2 AGN? What is the true connection between galaxy interactions, star formation and nuclear activity? What is the lifetime of these phenomena? How do LINERs fit in the general picture and are all of them AGN? Do different types imply different kind of objects, different evolutionary stages or a combination of evolution and line of sight orientation?  \\subsection{Triggering an AGN evolutionary sequence?}  Despite observational difficulties and limitations, there have been many attempts, based on different diagnostics, to investigate the possible triggering mechanisms of nuclear activity. Most agree that the accretion of material into a MBH (Lynden-Bell 1969) is the mechanism responsible for the emission, but how does the feeding of the black hole works? There are a lot of controversial results concerning this issue. It is known and widely accepted that interactions between galaxies can drive gas and molecular clouds towards its nucleus initiating an enhanced star formation (e.g. Li et al. 2008, Ellison et al. 2008). Many also believe that the same mechanism could give birth to an active nucleus (e.g. Umemura 1998, Kawakatu et al. 2006)  Indeed, there are several studies that conclude that there is an evolutionary sequence from starburst to Seyfert galaxies (e.g. Storchi-Bergmann et al. 2001;). Furthermore, based on the number and proximity of close ($\\mincir 60-100 \\;h^{-1}$ kpc) neighbours around the different type of active (Sy1, Sy2 and BIRG) galaxies ( e.g. Dultzin-Hacyan et al. 1999; Krongold et al. 2002; Koulouridis et al. 2006a,b) a very interesting evolutionary sequence has been suggested, starting with an interaction that triggers the formation of a nuclear starburst, which then evolves to a type 2 Seyfert, and finally to a Sy1. Such connection is likely independent of luminosity, as similar trends have been proposed for LINERs (Krongold et al. 2003) and ULIRGs and Quasars (Fiore et al. 2008 and references therein). This sequence can fully explain the excess of starbursts and type 2 AGN in interacting systems, as well as the lack of type 1 AGN in compact groups of galaxies (Martinez et al. 2008).  In addition, post starburst stellar populations have been observed around AGN (Dultzin-Hacyan \\& Benitez 1994; Maiolino \\& Rieke 1995; Nelson \\& Whittle 1996; Hunt et al. 1997; Maiolino et al. 1997; Cid Fernandes, Storchi-Bergmann \\& Schmitt 1998; Boisson et al. 2000, 2004; Cid Fernandes et al. 2001, 2004, 2005) and in close proximity to the core ($\\sim$50 pc), implying the continuity of these two states and a delay of 0.05-0.1 Gyr (M{\\\"u}ller S{\\'a}nchez et al. 2008) up to 0.5-0.7 Gyr (Kaviraj et al. 2008) between the onset of the starburst and the lighting up of the AGN. Davies et al.(2007), analyzing star formation in the nuclei of nine Seyfert galaxies found recent, but no longer active, starbursts which occurred 10 - 300 Myr ago. Furthermore, most of these studies (e.g. Hunt et al. 1997; Maiolino et al. 1997, Gu et al. 2001) separate type I from type II objects implying that recent star-formation is only present in type II objects. Support for the interactions-activity relation was recently provided by HI observations of Tang et al. (2008), who found that 94\\% of Seyfert galaxies in their sample were disturbed in contrast to their control sample (where only 19\\% were disturbed). We point out that a great theoretical success of the starburst/AGN connection is the quenching of the star formation induced by the AGN feedback which can explain the formation of red and dead elliptical galaxies (e.g. Springel et al. 2005a; Di matteo et al. 2005; Khalatyan et al. 2008). Recent results have found that AGN ionized outflows may carry enough energy to cease star formation in the host galaxy (Krongold et al. 2007, 2009; Chartas et al. 2008; Blustin et al. 2008). Furthermore, the correlation of morphologies of the host galaxies with the nuclear activity type can also lead us to a possible AGN evolutionary sequence (Mart{\\'{\\i}}nez et al. 2008).  There is also evidence pointing in the opposite direction. Li et al. (2008), analyzing SDSS data, showed that there is an increase in the starformation rate but no AGN enhancement associated to the presence of a close companion. However, they left a window open for the possibility of a delay between the onset time of the two phenomena. In any case it appears very likely that a non axi-symmetric potential, such as a tidal interaction, provides a triggering mechanism of activity.  In addition, the quenching of star formation could be due to AGN feedback, activated with a possible delay of $\\sim$0.5 Gyr, inducing an evolution from a Sy2 to a Sy1 state. Should this be true, the effect of interactions on the host galaxy would not be equally observable at the different phases of the AGN evolution. Of course, such an evolutionary scenario poses some problems to the simplest version of the unified scheme.  Observations of polarized broad emission lines from type 2 AGN galaxies gave strength to the unified picture, implying that this is scattered light by the torus. However, hidden broad lines are only observable in 50\\% of Seyfert 2s (Tran 2003) suggesting the possible existence of true type 2 objects, i.e. objects completely lacking a BLR. The lack of a BLR can be explained in terms of a very low accretion rate (Nicastro 2000, Elitzur 2008), and/or due to even greater obscuration (Shu et al. 2008). Supporting the low accretion model, we point out that unobscured type 2 AGN galaxies have also been observed by Brightman et al. (2008) and Bianchi et al. (2008).  We stress that the evolutionary scenario does not contradict the unification scheme. It implies that different types of AGN galaxies are in fact the same object (as the unification model proposes) but not necessarily at the same evolutionary phase. However, there could be a phase where only orientation defines the appearance of a Sy1 or a Sy2 state, which is the stage where the obscuring molecular clouds form a torus but have not yet been swept away, in agreement with the presence of polarized broad emission lines.  Taking into account the previous, when working with AGN samples, it is of crucial importance that the different types are analysed separately, in order to emphasize the similarities and the differences of each AGN type and not to smooth them by stacking them all together.   \\subsection{This work} This paper is the third in a series of 3-dimensional studies of the environment of active galaxies (Koulouridis et al. 2006a,b), based  on previous 2D analyses (Dultzin et al. 1999, Krongold et al. 2002), and will attempt to shed some more light to the starburst/AGN connection and to the evolutionary scenario proposed in our previous papers. It is a follow-up spectroscopic study aiming at investigating the effects of interactions on the neighbours of our Seyfert and Bright IRAS galaxies.  We will discuss our galaxy samples in section \\S 2. In \\S 3 we present the observations and data reduction. The spectroscopic analysis and classification of the galaxies is presented in \\S 4 and \\S 5, while the analysis of the available X-ray observations is given in \\S 6. Finally in sections \\S 7 and \\S 8 we will interpret our results and draw our conclusions. Due to the fact that all our samples are local, cosmological corrections of galaxy distances are negligible. Throughout our paper we use $H_{\\circ}=100 \\; h^{-1}$ Mpc. We use the term ``active'' for all galaxies that exhibit emission lines in their spectra, their origin being nuclear activity or starformation, since we believe that these two phenomena are strongly related to interactions. ",
        "conclusions": "We have found that the large majority ($>70\\%$) of close neighbours of AGN (being Sy1 or Sy2) and of Bright IRAS galaxies show activity, mostly starbursting but AGN as well. Furthermore, the close neighbours of Sy1 galaxies, being starbursts or AGN, are less ionised and thus seem to be a different, more evolved, population than those of Sy2s.  These results are in the same direction as those of our previous papers, ie., the significant excess of neighbours around Sy2s and BIRGs but not around Sy1s, and support an evolutionary sequence of galaxy activity, induced by interactions, % the main path of which is from inactivity, starburst, Sy2, Sy1 and back to inactivity, since the ``active'' phase should cease with AGN feedback. Finally, such sequence could be repeated by a new encounter.  The time needed for type 1 activity to appear should be larger than the timescale for an unbound companion to escape from the close environment, or comparable to the timescale needed for an evolved merger ($\\sim 1$Gyr, see Krongold et al. 2002). If our scenario is correct, then there should be two distinct populations of Sy2s: the obscured ones which pre-exist the Sy1 phase and should exhibit close companions and the ``naked'' ones which appear later and should be isolated.  The results of our analysis also provide valuable guidelines for future studies of the AGN environment, since they point out two very important issues:  \\noindent [a.] Taking under consideration the possible evolutionary scheme discussed and if indeed strong interactions can be traced only until the first Sy2 phase of the phenomenon, the use of large AGN samples, without discriminating the different activity types, would introduce a large confusion to the interpretation of the results of studies of the AGN environment. For example, keeping indiscriminately the huge population of LINERs in analyses of AGN samples could result in confusion regarding the AGN local environment.  \\noindent [b.] Starburst galaxies should not be handled under any circumstances as normal galaxies, since the starburst/AGN connection appears quite plausible and composite objects do exist anyway. In a nutshell, all different categories of active galaxies should be treated separately.  There are still many unresolved issues and caveats concerning these studies, since the evolutionary sequence is not unique and should also depend on the geometry, the density and other factors of the obscuring and the accreting material, as well as on the mass of the host galaxy and its black hole."
    },
    "0910/0910.0849_arXiv.txt": {
        "abstract": "If the cosmological inflationary scenario took place in the cosmic landscape in string theory, the inflaton, the scalar mode responsible for inflation, would have meandered in a complicated multi-dimensional potential. We show that this meandering property naturally leads to many e-folds of inflation, a necessary condition for a successful inflationary scenario. This behavior also leads to fluctuations in the primordial power spectrum of the cosmic microwave background radiation, which may be detected in a near future cosmic variance limited experiment like PLANCK. ",
        "introduction": "It is generally believed that there is a cosmic landscape in string theory \\cite{KKLT}. One expects the string scale to be in between the Hubble and the Planck scale, $H \\ll m_s \\ll \\mpl$, so that a single Planck scale field range in the landscape naturally contains many string scale features. When the vacuum energy is not too far below the string scale, the universe is expected to undergo cosmic inflation, which describes how our universe began. The inflationary universe scenario is strongly supported by theoretical considerations and observational data. For our purpose here, the cosmic landscape may be approximated by a $d$-dimensional random potential, and the inflaton (the scalar mode for inflation) path appears smooth in the coarse-grained limit, but is actually rugged with fine-grained string scale features. In this paper, we will present two generic properties of such an inflationary scenario: \\begin{enumerate} \\item At least 60 e-folds of inflation is necessary in an inflationary scenario to solve the well-konwn cosmological problems such as the flatness and the horizon problems. To achieve this in the slow-roll inflationary scenario usually requires an almost flat smooth potential. Without any fundamental theory, one can always invent such a potential. However, in the context of string theory, such an almost flat potential requires some fine-tuning. While a generic potential may not be flat enough, it is not necessarily smooth either. A rugged $d$-dimensional potential is more realistic in the context of the cosmic landscape (Fig.\\ref{fig_pot}), where $d \\gg 1$. For large enough $d$, percolation probability approahes unity so the wave function of the universe is mobile in the landscape \\cite{Tye:2006tg,Sarangi:2007jb}. For a Planck scale field range, one expects many string scale features and consequently the inflaton takes a meandering path. Since the de-Sitter quantum fluctuation ($H$) is smaller than the size of the features ($m_s$), the turns of the inflaton path may be treated as classical scatterings. We will prove that an arbitrary detour of the inflaton path always increases the travel time, leading to more e-folds than given by the corresponding coarse-grained smooth potential. With appropriately many features in the potential, enough e-folds of inflation becomes generic. Although the scenario is motivated by the cosmic landscape, if a specific string compactification is complicated enough, a brane motion inside the bulk can exhibit similar meandering behavior.  \\item The above scenario naturally leads to fluctuations in the the cosmic microwave background radiation (CMB) temperature  (TT) and polarization  (TE/EE) power spectra. These fluctuations are due to the turns and detours of the inflaton path in the $d$-dimensinal potential. Due to the randomness of the features in the potential, such fluctuations appear with irregular spacings, magnitudes and shapes \\cite{Tye:2008ef}. While some of these features are probably too fine and/or small to be detected, a few big fluctuations may already have been observed in the  data \\cite{Nolta:2008ih,Nagata:2008tk, Nagata:2008zj,Nicholson:2009pi}. With more data coming in the future, one expects this prediction to be well tested. The existence of such fluctuations will have a profound impact on the inflationary scenario, on the origin of our universe and on the existence and properties of the cosmic landscape. \\end{enumerate} ",
        "conclusions": "In the paper, we have presented a scenario where a large number of scalar fields collectively contribute to cosmic inflation. It is important to distinguish our scenario from existing multifield scenarios like assisted inflation \\cite{Liddle:1998jc} and N-flation \\cite{Dimopoulos:2005ac, Easther:2005zr}.  In the assisted inflation scenario, different fields have identical uncoupled potential. The trajectory of inflaton field is most likely a smooth path close to a straight line during the few e-folds relevant for CMB observation, which does not allow entropic perturbations to play any role. If the single field potential has some features like steps or kinks, each field will experience the same feature as it rolls down the potential. This will create identical features on the power spectrum with the same shape/magnitude.  In the N-flation scenario, a large number of axion fields play the role of the inflatons. If we only keep the leading order one-instanton terms, and Taylor expand each axion potential around its minimum, we get a set of fields with different masses. At one-instanton level, different inflaton fields still decouple, but since their masses are all different, the inflaton trajectory could be curved. However, the turning of the inflaton field is very mild in the slow-roll regime, and again does not lead to sharp features in the power spectrum. One can keep the sub-leading multi-instanton terms in the potential; this will explicitly introduce couplings between the fields, and could make the inflaton path more convoluted, therefore providing an explicit realization of the meandering inflaton in string theory. In this case, our approach provides a different perspective to the analysis of these complicated scenarios. In particular, our analysis implies that the simplified analysis carried out in the literature on such a scenario may substantially under-estimate the e-folds produced by the model.  There has been a lot of conjectures regarding the features in the CMB power spectrum. However, it is not clear what is the origin of such sharp features from string theory. In Ref.\\cite{Bean:2008na} and \\cite{Flauger:2009ab}, features in the potential are well motivated to arise from duality cascade or monodromy. Here we provide another natural motivation: the meandering of the inflaton in the cosmic landscape. Regarding observations, duality cascade and monodromy lead to features on the power spectrum with predicted location and magnitude, while here we have random fluctuations in the power spectrum which is extremely hard to be picked out by current data analysis approaches.   \\vspace{0.3cm} {\\bf Acknowledgments:} We thank Rachel Bean, Scott Dodelson, Liam McAllister, David Spergel, Jun'ichi Yokoyama, Yang Zhang for valuable discussions. This work is supported by the National Science Foundation under grant PHY-0355005."
    },
    "0910/0910.0552_arXiv.txt": {
        "abstract": "NLTE radiative transfer calculations of differentially expanding supernovae atmospheres are computationally intensive and are almost universally performed in time-independent snapshot mode, where both the radiative transfer problem and the rate equations are solved assuming the steady-state approximation. The validity of the steady-state approximation in the rate equations has recently been questioned for Type II supernova (SN II) atmospheres after maximum light onto the plateau.  We calculate the effective recombination time of hydrogen in SN~II using our general purpose model atmosphere code \\texttt{PHOENIX}.  While we find that the recombination time for the conditions of SNe~II at early times is increased over the classical value for the case of a simple hydrogen model atom with energy levels corresponding to just the first 2 principle quantum numbers, the classical value of the recombination time is recovered in the case of a multi-level hydrogen atom.  We also find that the recombination time at most optical depths is smaller in the case of a multi-level atom than for a simple two-level hydrogen atom. We find that time dependence in the rate equations is important in the early epochs of a supernova's lifetime. The changes due to the time dependent rate equation (at constant input luminosity) are manifested in physical parameters such as the level populations which directly affects the spectra. The $H_\\alpha$ profile is affected by the time dependent rate equations at early times. At later times time dependence does not significantly modify the level populations  and therefore the $H_\\alpha$ profile is roughly independent of whether the steady-state or time-dependent approach is used. ",
        "introduction": "Supernovae are one of the most widely researched objects in astrophysics.  There has been much debate about the detailed physical nature of these objects, driven by the discovery of the dark energy using SNe Ia \\citep{riess_scoop98,perletal99}. Supernovae add rich variety to the metal abundances in galaxies and are important to star formation theories. Supernova nucleosynthesis of heavy elements is responsible for galactic chemical evolution.  Prior to the discovery of dark energy, supernova research was spurred on by the discovery of the peculiar Type II SN 1987A in the LMC \\citep[see][and references therein]{sn87arev89}.  Type II supernovae (SNe II) are classified by the presence of strong Balmer lines in their spectra in contrast with SNe I which lack strong Balmer lines.  SNe II are due to core collapse of massive stars.  SN~1987A was extremely well observed and much work has been done to explain all aspects of its spectra. The Balmer lines were not well reproduced at early times (especially about 8 days after explosion).  At first very large overabundances of the s-process elements Ba and Sc were suggested \\citep{williams87A87,williams87A87b,hof87A88,danziger87A88}. Further work showed that the Ba~II and Sc~II lines could be  reproduced by assuming that the s-process abundances were enhanced by a factor of 5 compared to the LMC abundances \\citep{mazz87A92}.  The  hydrogen Balmer line problem in many SNe~II has not been accounted in a consistent manner.  \\citet{dessart08} summarizing their work using time-independent rate equations found that they were unable to reproduce the $H_\\alpha$ line after 4~days in SN 1987A, 40~days in SN~1999em, and 20~days in SN~1999br. \\citet{Chugai} point out the lack of a consistent physically motivated framework to explain the early epoch Balmer profile that matches with the observed line strength. Other authors followed different approaches to reproduce the observed line strengths with varying degrees of success \\citep{h03a,schm90,HM98,dessart05b}. \\citet{Chugai} argued that the poor fit of the Balmer lines could be overcome by including time-dependence into the hydrogen rate equations.  They argued that even in the intermediate layers, the recombination time increases due to Lyman alpha trapping or ionization from the first excited state.  Consequently, the recombination time becomes comparable to the age of the supernova, making the steady-state approximation suspect.  We numerically estimated the recombination time by directly calculating the recombination rate from the continuum into the bound states of hydrogen using H in NLTE with 31 bound states and all other metals in LTE.  We use the general purpose model atmosphere code \\texttt{PHOENIX} developed by \\citet{hbjcam99}.  We specify the density structure $\\rho \\propto r^{-7}$ and assume the ejecta mass to be 14.7~\\msun\\ as is appropriate for SN~1987A \\citep{nomoto1988}.  \\texttt{PHOENIX} calculates the level populations, temperature, and radiation field for successive epochs with a time interval of approximately two days.  We generated our models and spectra using both time-dependent and time-independent rate equations.  We calculated our recombination time for each epoch from the best fit spectra in each case.  In \\S~2, we outline the theoretical framework for the recombination time and motivation for our work.  In \\S~3, we discuss earlier work done to test the relevance of time-dependence in type II supernova atmospheres. In \\S~4, we describe the basic framework for our code and in \\S~5 we present our approach to compute the recombination time. In \\SS~6--8 we analyze our results and address the question of the necessity of incorporating time dependence in the rate equations due to the increased recombination time. ",
        "conclusions": "We have studied the importance of time-dependence in the rate equations in a type II supernova atmosphere due to the increase in the recombination time. We also explored the effects of having many angular momentum sub-states. In our description we use $\\tau_{H(4)}$ and $\\tau_{H(31)}$ as the recombination time for a 4-level hydrogen atom model and 31-level hydrogen atom model respectively.  1) \\uc05 proposed that $\\tau_{H(4)}$ is of the order of $10^{7}$~s at $T \\sim 5000$~K and $n_{e}=10^{8}$ cm$^{-3}$. For epochs later than 6 days, we calculated $\\tau_{H(4)}$ at similar electron density and temperature to be around $10^{7}$~s ($\\taustd \\sim 0.02$).  Thus, we recover the recombination time scale predicted by \\uc05 for the 4-level hydrogen case. An important difference between our results and those of \\uc05 is that  we find the electron densities and temperatures which they used as photospheric conditions to occur at $\\taustd \\sim 0.02$ and not near $\\taustd = 1.0$. Around $\\taustd=1.0$, the free electron density even for Case 1 is much higher and hence the recombination time is much smaller. This optical depth mismatch is due to the fact that we are looking at the continuum optical depth and not the line optical depth. The ratio of the Balmer line opacity to the continuum opacity  is $\\sim 10^{3}$ in the recombining region.    We also compare the classical recombination time (calculated using $1/\\alpha n_{e}$)  with our calculated $\\tau_{H(4)}$ generated from the solution of the radiative transfer equation. We refer to the approximate classical recombination time as $\\tau_{anal}$. This allows us to determine the factor $w_{21}$ (\\uc05). Figure~\\ref{tanal_by_th} displays the ratio $\\tau_{anal}/\\tau_{H(4)}$,  a direct measure of $w_{21}$ described by \\uc05.   In Figure~\\ref{ratio_recombination_time}, we display the ratio of the recombination time for Case 2 to the recombination time from a simple  hydrogen  atom (Case 1)  at the same epoch and at the same optical depth.  1) For all epochs, $\\tau_{H(31)}$ is smaller than $\\tau_{H(4)}$ at almost all optical depths. For higher optical depths, the ratio increases and approaches unity.  because the ionization fraction for hydrogen and $\\tau_{H(31)}$ becomes independent of the number of bound levels.  At this point the system is almost in LTE and the number of free electrons depends only on the temperature.  2) Figure \\ref{tanal_by_th} shows that the simple hydrogen atom model (4-level) overestimates the recombination time by a factor of 100 or more at all almost all epochs and optical depths.  3) Figure \\ref{tanal_by_th31} shows that the recombination time recovered from a multi-level hydrogen atom compares with the standard value ($\\tau_{anal}$) for $\\tau_{std}<1.0$. At the ionization front $\\tau_{H(31)}$ differs significantly from $\\tau_{anal}$. 4) Based on our work on SN 1999em, we find that the transient nature of the level population of hydrogen is not a crucial factor on the plateau, but it is important in the active recombining regime. This effect could be important for early features in the spectra but the steady state approximation should be valid in the plateau phase.  5) In the early early spectra of SN~1987A we find that $H_\\alpha$ is better fit using the time-dependent formulation.  6) The recombination time jumps above the underlying trend only at the ionization front. This is just caused by the sudden reduction of the electron density driving the the recombination time higher in the ionization front.    It is very difficult to decouple the effects of time dependent rate equations in a radiative transfer framework and the effects of multi-level model atoms by only looking at spectra. The Balmer profile of the spectra of SN 1987A at early times (Figure \\ref{spectra_observed}) and the level populations from our results on SN 1999em (Figure \\ref{bi_99em_n1} -- \\ref{bi_99em_n3}), indicate the effects of the time dependent rate equations at early times. The increased recombination time (Figure \\ref{recombination_time}) at early times for SN 1987A also supports the importance of time-dependent rate equations. Therefore we conclude that time dependence is more important at early times than later times. The effect of multi-level atoms can be seen in Figure \\ref{recombination_time} which clearly shows than even at later times $\\tau_{H(4)}$ is much higher than $\\tau_{H(31)}$. Therefore conclusions using only a 4-level model may overestimate the importance of time dependence. Of course, including time-dependence is not terribly costly numerically and thus it can be included when a time sequence is calculated."
    },
    "0910/0910.5007_arXiv.txt": {
        "abstract": "Baryon acoustic oscillations (BAO) provide a robust standard ruler with which to measure the acceleration of the Universe. The BAO feature has so far been detected in optical galaxy surveys. Intensity mapping of neutral hydrogen emission with a ground-based radio telescope provides another promising window for measuring BAO at redshifts of order unity for relatively low cost. While the cylindrical radio telescope (CRT) proposed for these measurements will have excellent redshift resolution, it will suffer from poor angular resolution (arcminutes at best). We investigate the effect of angular resolution on the standard ruler test with BAO, using the Dark Energy Task Force Figure of Merit as a benchmark. We then extend the analysis to include variations in the parameters characterizing the telescope and the underlying physics. Finally, we optimize the survey parameters (holding total cost fixed) and present an example of a CRT BAO survey that is competitive with Stage III dark energy experiments. The tools developed here form the backbone of a publicly available code that can be used to obtain estimates of cost and Figure of Merit for any set of survey parameters. ",
        "introduction": "A standard ruler test with Baryon acoustic oscillations (BAO) is considered the most robust and systematics-free method to probe the dark energy equation of state \\citep{Albrecht:2006um}. The sound waves which propagated through a mixture of photons and baryons in the early Universe left a distinct oscillatory signature in the Cosmic Microwave Background (CMB) and the large scale structure of matter (or galaxies) with a characteristic scale set by the sound horizon at the epoch of recombination \\citep{Peebles70,Holtzman89,Hu96,EH98}. Measurements of the CMB provide this sound horizon scale, and we can then use the BAO in low-redshift clustering as a standard ruler to measure the distances to various redshifts, and therefore the equation of state of dark energy \\citep{HW96,EHT98,Eisen03,Blake03,Linder03,Hu03,SE03}.  While detections of BAO so far have used the optical band of the electromagnetic spectrum, using spectroscopic or photometric techniques, there has been growing interest in the feasibility of using 21cm emission from neutral hydrogen \\citep{WL08,LW08,WL09,Visbal09} to detect BAO at high \\citep{Mao:2008ug,WLal08} and medium  \\citep{Chang:2007xk,Ansari:2008yw} redshifts. Observing BAO from intensity mapping of 21cm emission, especially at medium redshift near $z\\sim 1$, has several advantageous features. The BAO is a relatively weak feature on large scales, and therefore very large survey volumes are needed. A ground-based cylindrical radio telescope (CRT) with a large field of view can easily cover most of the sky in a relatively short time. Second, the electronics required for frequencies near $\\nu\\sim 1{\\rm GHz}$ are cheap and easy to build. Digital electronics with precise timing offer high precision (better than ppm) frequency and, hence, redshift measurements. Third, radio surveys rely on different tracers of large scale structure (neutral hydrogen) than do optical survey (luminous galaxies). Seeing the signal in two different sets of mass tracers would be compelling evidence for its robustness.  A disadvantage of radio waves is their longer wavelengths, making it more difficult to obtain high angular resolution. For example, a 100m by 100m telescope will have an angular resolution of $(\\lambda/L) \\sim 15$ arcminutes at $z\\sim 1$. Radio observations of cosmological hydrogen are severely contaminated by foreground sources, primarily Galactic synchrotron radiation. We will not consider the problem of foreground removal here, but will show that foregrounds are smooth at the BAO scale and then assume that they can be eliminated to the accuracy of the statistical errors.  In this paper, we consider a ground-based 21cm intensity mapping with a CRT targeting a redshift range of 0.2-2, motivated by the study of \\citet{Chang:2007xk}. In \\S~\\ref{sec:StoN}, we derive the signal-to-noise of the power spectrum for a given set of radio telescope configuration parameters, discuss galactic shot noise and the effect of foregrounds, and explain the assumptions used in our method of Fisher matrix projections. We model the effects of angular resolution with a window function that approximates the effects of a typical CRT telescope response function on the power spectrum.  Using both Fisher matrix analysis and Monte Carlos, we study the issue of angular resolution in detail in \\S~\\ref{sec:antheta}, focusing on the impact of resolution on the determination of dark energy parameters. Then in \\S~\\ref{sec:FoM}, we consider a simple compact CRT telescope, and investigate the Dark Energy Task Force (DETF) Figure of Merit (FoM) \\citep{Albrecht:2006um} as a function of telescope parameters. The main result of the paper is captured in Fig.~\\ref{fig:figparam} where the FoM is plotted as a function of these parameters. In \\S~\\ref{sec:Max}, we optimize the telescope configuration parameters for the maximum FoM and show that an inexpensive CRT BAO survey can be competitive with Stage III dark energy experiments \\citep[for optimazation studies in optical surveys, see][]{Parkinson07,Parkinson09}. In \\S~\\ref{sec:con} we summarize the results. ",
        "conclusions": ""
    },
    "0910/0910.4830_arXiv.txt": {
        "abstract": "We present the results of some experiments performed in the Padua simulators of planetary environments, named LISA, used to study the limit of bacterial life on the planet Mars. The survival of {\\it Bacillus} strains for some hours in Martian environment is shortly discussed. ",
        "introduction": "Planet Mars is a typical target for exploring the possibility that life arose in a planet different than Earth. Its present surface conditions are very inhospitables, if compared with most of the terrestrial environments, but several lifeforms have shown a strong capability to adapt to very harsh conditions and to survive even for some period in circumterrestrial space. If life, similar or different from the terrestrial one, existed on Mars in the past, some  life forms could have adapted to the climate changes and could have survived in some ecologic niche, near the soil or deep under the surface. Moreover, while exploring Mars, we may incidentally or voluntary drop terrestrial lifeforms, able to survive and to be reactivated once brought back in laboratories. Finding conditions that allow the survival of lifeforms in a Martian environment can have a double value: increasing the hope of finding extraterrestrial life and defining the limits for a terrestrial contamination of planet Mars. ",
        "conclusions": "In our experiments, we found that desiccation  effect (water escape because of low pressure) may strongly decrease the survival of vegetative cells, but not of spores. In Martian environment, UV light appears to have the most cytocidal effect, while atmospheric gases or temperature are not relevant to the survival of cells or spores. Vegetative cells are inactivated by UV light in a few minutes, while spores may survive for hours (Fig. \\ref{fig:fig1}). Two of our {\\it Bacillus} strains, {\\it B. pumilus} and {\\it B. Nealsonii}, have a particular capability to survive in Martian conditions without being screened by dust or other shields. As endospores suspension, they survive at least 4 hours and in some cases up to 28 hours in Martian conditions.  We simulated the dust coverage happening on the real planet by blowing on the samples a very smaller quantity of grain of volcanic ash or dust of red iron oxide. Samples covered by these dust grains have shown a high percentage of survival, indicating that under the surface dust, if life was present on Mars in the past, some bacteria cell could still be present.   \\begin{figure} \\resizebox{13.5cm}{!}{\\includegraphics{Fig_1.eps}} \\caption{Survival of spores in LISA simulator at  +23 and -80 $^\\circ$C vs. exposure to UV light.}\\label{fig:fig1} \\end{figure}"
    },
    "0910/0910.3767_arXiv.txt": {
        "abstract": "Abundance anomalies observed in globular cluster stars indicate pollution with material processed by hydrogen burning.  Two main sources have been suggested: asymptotic giant branch (AGB) stars and massive stars rotating near the break-up limit (spin stars). We discuss the idea put forward by \\citet{DeMink+09b} that massive binaries may provide an interesting alternative source of processed material.  We discuss observational evidence for mass shedding from interacting binaries.  In contrast to the fast, radiatively driven winds of massive stars, this material is typically ejected with low velocity. We expect that it remains inside the potential well of a globular cluster and becomes available for the formation or pollution of a second generation of stars. We estimate that the amount of processed low-velocity material that can be ejected by massive binaries is larger than the contribution of two previously suggested sources combined. ",
        "introduction": "For a long time star clusters have been considered as idealized single-age, chemically homogeneous stellar populations. However, it has recently become clear that some clusters show multiple main sequences and sub-giant branches and extended horizontal branches \\citep[e.g.][]{Piotto+07}.  The presence of these features in some intermediate-age clusters may be explained by the effects of stellar rotation on the location of stars in the colour magnitude diagram \\citep{Bastian+DeMink09}. However, this provides no explanation for the features seen the old globular clusters, which seems to imply the existence of multiple populations within one cluster.  In addition, large star-to-star abundance variations are found for light elements such as C, N, O, Na and Al, while the composition of heavier elements (Fe-group and $\\alpha$-elements) seems to be constant.  Field stars with the same metallicity do not exhibit these abundance patterns \\citep[for a review see][]{Gratton+04}. These chemical variations have been interpreted as originating from the presence of both a ``normal'' stellar population, exhibiting abundances similar to field stars of the same metallicity, and a second  population of stars formed out of material processed by hydrogen burning via the CNO-cycle and the NeNa and MgAl chains \\citep[e.g.][]{Prantzos+07}.  {According to \\citet{Carretta+09},  50-70\\% of the stars in gloular clusters belong to the second population.}  Two sources of processed ejecta have been proposed: the slow winds of \\emph{massive AGB stars}, which enrich their convective envelopes with H-burning products \\citep{Ventura+01,Dantona+02,Denissenkov+Herwig03} and fast-rotating massive stars (we refer to these as \\emph{spin stars}), which are believed to expel processed material centrifugally when they reach break-up rotation \\citep{Prantzos+Charbonnel06,Decressin+07a}.% In this scenario a first generation of stars is formed out of pristine material. Their low-velocity ejecta are trapped inside the potential well of the cluster and provide the material for the formation of a second generation of stars% \\footnote{\\citet{Glebbeek+09} suggested that a chain of multiple stellar collisions in the dense center of a star cluster may also enrich the interstellar medium with processed material.}.  Although both proposed sources are promising, matching the observed abundance patterns and providing enough ejecta for the formation of a second generation which outnumbers the first generation have proven to be two major challenges. In this Letter we propose \\emph{massive binaries} as a candidate for the internal pollution of globular clusters.     \\begin{figure}[] \\vspace*{-2.0 cm} \\begin{center} \\includegraphics[width=\\textwidth]{cartoon.eps} \\caption{Cartoon representation of the proposed scenario. {\\bf Left}: A binary system at the onset of mass transfer. The deepest layers in the donor star have been processed by proton capture reactions. The accreting star spins up as it accretes mass and angular momentum until it approaches the break-up limit. {\\bf Right}: The same system after the donor star has been stripped from its envelope. The companion star acrreted just a fraction of the transferred mass, mainly unprocessed material originating from the outermost layers of the donor star.  Material orginating from deeper layers in the donor star are shedded into a circumbinary disk.    } \\label{fig1} \\end{center} \\end{figure} ",
        "conclusions": "We discussed the potential of massive binaries as the source of enrichment in globular clusters. The majority of massive stars are expected to be members of interacting binary systems. These return most of the envelope of their primary star to the interstellar medium during non-conservative mass transfer.  We show that the amount of polluted material ejected by binaries may be larger than that of the two previously suggested sources: massive AGB stars and the slow winds of fast rotating massive stars. After dilution with pristine material, as lithium observations suggest, binaries could return enough material for the formation of a chemically enriched second generation that is equally numerous as the first generation of low-mass stars, without the need to assume a highly anomalous IMF, external polution of the cluster or a significant loss of stars from the non-enriched first generation.  In addition to providing a new source of slowly-ejected enriched material, binary interaction also affects the previously proposed scenarios.  Binary mass transfer naturally produces a large number of fast-rotating massive stars which may enrich their surroundings even further.  Binary interaction will also affect the yields of intermediate-mass stars. Premature ejection of the envelope in 4-9\\Msun~stars will result in ejecta with less pronounced anti-correlations as suggested in the AGB scenario. On the other hand, we expect that binary-induced mass loss may also prevent the dredge-up of helium burning products.  For a detailed comparison of the chemical predictions of this scenario binary models {for a range of masses and orbital periods} are needed and population synthesis models are essential to fullly evaluate the mass budget of the different sources.  {Some peculiar feature, such as the apparant presence of distinct, chemically homogeneous subpopulations in  $\\omega$~Cen and NGC~2808 \\citep[e.g.][]{Renzini08} deserves further attention.}"
    },
    "0910/0910.1248_arXiv.txt": {
        "abstract": "We present the luminosity function and selection function of 60 $\\mu m$ galaxies selected from the Imperial IRAS-FSC Redshift Catalogue (IIFSCz). Three methods, including the $1/V_{\\textrm{max}}$ and the parametric and non-parametric maximum likelihood estimator, are used and results agree well with each other. A density evolution $\\propto (1+z)^{3.4\\pm0.9}$ or a luminosity evolution $\\propto \\exp(1.7~t_L / \\tau)$ where $t_L$ is the look-back time is detected in the full sample in the redshift range [0.02, 0.1], consistent with previous analyses. Of the four infrared subpopulations, cirrus-type galaxies and M82-type starbursts show similar evolutionary trends, galaxies with significant AGN contributions show stronger positive evolution and Arp 220-type starbursts exhibit strong negative evolution. The dominant subpopulation changes from cirrus-type galaxies to M82-type starbursts at $\\log_{10} (L_{60} / L_{\\odot}) \\approx 10.3$.  In the second half of the paper, we derive the projected two-point spatial correlation function for galaxies of different infrared template type. The mean relative bias between cirrus-type galaxies and M82-type starbursts, which correspond to quiescent galaxies with optically thin interstellar dust and actively star-forming galaxies respectively, is calculated to be $b_{\\textrm{cirrus}}/b_{\\textrm{M82}}=1.25\\pm0.07$. The relation between current star formation rate (SFR) in star-forming galaxies and environment is investigated by looking at the the dependence of clustering on infrared luminosity. We found that M82-type actively star-forming galaxies show stronger clustering as infrared luminosity / SFR increases. The correlation between clustering strength and SFR in the local Universe seems to echo the basic trend seen in star-forming galaxies in the Great Observatories Origins Deep Survey (GOODS) fields (at $z\\sim1$). ",
        "introduction": "The study of galaxy distribution on large scales ($\\sim10 ~h^{-1}$ Mpc) where density fluctuations are still in the linear regime ($\\delta \\rho / \\rho \\ll 1$) has been an important tool in constraining cosmological models. However, the unknown relationship between the distribution of dark matter and that of luminous objects, usually referred to as galaxy bias, has complicated the extraction of cosmological parameters, especially when we compare results from surveys of different galaxy subpopulations or of varying depths. At small scales ($\\sim1 ~h^{-1}$ Mpc) where nonlinear effects become important, detailed studies of galaxy bias can improve our understanding of galaxy formation and evolution which is still unclear due to the convolution of gravitational evolution and complex astrophysical processes (e.g. gas cooling, fragmentation, star formation, feedback). This paper attempts to measure the relative galaxy bias over the intermediate scale ($1\\sim10$) Mpc, i.e. the weakly nonlinear scale, as we are constrained by low galaxy density at small scales.  Several models of galaxy bias have been developed over the past 20 years or so. The simplest one is to assume galaxies trace the underlying mass in a linear, deterministic way, $\\delta_{\\textrm{g}} = b \\delta_{\\textrm{m}}$ where $\\delta_{\\textrm{g}}$ and $\\delta_{\\textrm{m}}$ is the galaxy and matter density perturbation (or contrast) field and $b$ is the linear bias parameter. In the peaks bias model (Kaiser 1984; Bardeen et al. 1986), the relationship between the correlation function of galaxies and that of the underlying dark matter satisfies $\\xi_{\\textrm{g}}(r)=b^2 \\xi_{\\textrm{m}}(r)$, based on statistics of the density peaks in a Gaussian random field. On large scales, it seems reasonable to assume a scale independent linear galaxy bias if galaxies form in virialized dark matter halos. On intermediate and small scales, the situation is more complicated. In more realistic galaxy bias models, $b$ is thought to be time-varying, scale-dependent, non-local, or even stochastic to some extent. N-body simulations of dark matter and semi-analytic models of galaxy formation predict a scale-dependent relation between galaxies and dark matter due to the mass-dependence of galaxy formation efficiency as well as a luminosity-dependent bias for galaxies with luminosities above $L^*$ (Benson et al. 2000; Benson et al. 2001). Attempts to understand galaxy bias have also been made in the halo occupation distribution (HOD) framework where the relationship between the galaxy distribution and the underlying dark matter is encoded in the conditional probability function $P(N|M)$ which gives the number of galaxies of a given type in a halo of virial mass $M$ (Jing et al. 1998; Ma \\& Fry 2000; Peacock \\& Smith 2000; Cooray \\& Sheth 2002; Zehavi et al. 2005; Tinker et al. 2005; Zheng, Coil \\& Zehavi 2007; Simon et al. 2008).  Observationally, several manifestations of galaxy bias are known. From dipole analysis, different values of the $\\beta \\equiv \\Omega_m^{0.6}/b$ parameter have been reported using different types of mass tracer (Rowan-Robinson et al. 2000; Kocevski, Mullis \\& Ebeling 2004; Erdo\\u{g}du et al. 2006 and references therein). The well-known morphology-density relation (Dressler 1980; Goto et al. 2003) states that late-type galaxies (disk-dominated) prefer low density regions and early-type (bulge-dominated) galaxies dominate in dense cluster cores. There is also a correlation between colour and morphology (the majority of blue galaxies are late-type galaxies, while red galaxies are mostly early-type galaxies), as well as a correlation between colour and luminosity (red galaxies are systematically more luminous than blue galaxies). Therefore, galaxy clustering is expected to be dependent on physical properties such as morphology, colour, surface brightness and luminosity. Using a total of over $100,000$ galaxies from the 2dF Galaxy Redshift Survey (2dFGRS), Norberg et al. (2001) found a scale-independent luminosity dependence of galaxy bias of the form $b/b^*=0.85+0.15L/L^*$ which means the luminosity-dependence is conspicuous above $L^*$. It agrees with the conclusions of Benoist et al. (1996) based upon 3,600 galaxies selected from the Southern Sky Redshift Survey 2 at the faint end, but it exhibits a less steep rise at the bright end. Zehavi et al. (2002, 2005) studied galaxy clustering in the SDSS, which not only confirmed the conclusions of Norberg et al. (2001) but also probed the relative bias behaviour to fainter galaxies. Madgwick et al. (2003) divided their sample from the 2dFGRS into relatively passive and active star-forming galaxy subsamples. A scale-dependent relative bias was found in comparing the clustering strength of the two subsamples. On scales smaller than 8 $h^{-1}$ Mpc, the relative bias between the two spectral classes is $b_{\\textrm{passive}}/b_{\\textrm{active}}=1.45\\pm0.14$. On scales larger than 10 $h^{-1}$ Mpc, the relative bias reduces to unity. Zehavi et al. (2002; 2005) used a luminosity-dependent colour-cut to divide their volume-limited sample from the SDSS. The colour-dependence of the correlation function is found to be very similar to that found in Madgwick et al. (2003). In addition, the luminosity-dependence for red galaxies is such that both faint and luminous red galaxies cluster more strongly than moderately bright red galaxies around $L*$ while a luminosity-dependence for blue galaxies is not noticeable (Wang et al. 2007; Swanson et al. 2008; Cresswell \\& Percival 2009).  The optical colour of a galaxy is connected with a complicated combination of star formation history and current star formation rate (SFR). In contrast, it is well known that the far-infrared emission is a good SFR indicator due to the re-processing of starlight by dust. So in principle one can measure clustering strength of infrared galaxies to learn about the relation between environment and SFR. So far, little is known about the clustering of infrared galaxies and its dependence on various physical parameters. The largest sample of local infrared galaxies is provided by IRAS. The correlation length of IRAS galaxies ($r_0\\sim4 ~h^{-1}$ Mpc) is similar to that of local star-forming/blue galaxies found in the SDSS and 2dFGRS (Saunders et al. 1992; Fisher et al. 1994). Clustering measurements of infrared galaxies did not show significant dependence on infrared luminosity (Beisbart \\& Kerscher 2000; Szapudi et al. 2000; Hawkins et al. 2001). This was thought to be a result of the combined effect of the dark matter halo mass and SFR in determining the far-infrared luminosity. From the SFR - density relation, we know that star-forming galaxies preferentially reside in low density environments (G\\'{o}mez et al. 2003). On the other hand, the gas content in a galaxy is largely dependent on the mass of the host dark matter halo. The dependence of clustering on infrared galaxy types is not entirely clear either. With a few thousand galaxies in the QDOT IRAS galaxy redshift survey (Lawrence et al. 1999), Mann, Saunders \\& Taylor (1996) did not detect any significant variation in the clustering of warm and cool IRAS galaxies separated by the far-infrared emission temperature (above or below 36 K respectively), assuming the far-infrared spectrum can be fitted by $S_{\\nu}\\propto \\nu B_{\\nu}(T)$. Hawkins et al. (2001), using the much bigger IRAS PSCz sample (Saunders et al. 2000), found that the relative bias between cool and warm infrared galaxies is around 1.5 on scales between 1 and 10 $h^{-1}$ Mpc.  Our aim is to investigate the relation between star formation activity in star-forming galaxies and environment using infrared samples in the general field in the local Universe. We study the clustering dependence on infrared luminosity and infrared template type, taking advantage of a new redshift catalogue of over 60,000 galaxies selected at $60~\\mu m$ from the IRAS Faint Source Catalog (FSC). The remainder of this paper is organised as follows. Section~\\ref{sample} briefly describes the galaxy sample used. Details of the new catalogue can be found in Wang \\& Rowan-Robinson (2009). In Section~\\ref{LF and SF}, we derive the luminosity function and the selection function using various methods which will then be used to construct random catalogues in measuring the clustering signal with flux-limited samples. In Section~\\ref{galaxybias}, clustering dependence on infrared template type and infrared luminosity is investigated using the projected two-point spatial correlation function. Finally, discussions and conclusions are presented in Section~\\ref{discussions and conclusions}. Throughout the paper, unless otherwise stated, we assume $H_0=72~\\textrm{km}~\\textrm{s}^{-1}~\\textrm{Mpc}^{-1}$ and a flat universe with $\\Omega_M=0.3$ and $\\Omega_{\\Lambda}=0.7$. Luminosities are quoted in units of $h_{72}^{-2}~L_{\\odot}$. ",
        "conclusions": "\\label{discussions and conclusions}  We have derived the luminosity function and selection function of the IIFSCz to learn about the evolution of infrared galaxies. To minimize the effect of incompleteness, we chose a flux limit of 0.36 Jy at 60 $\\mu$m. Either a pure density evolution $P=3.4$ or pure luminosity evolution $Q=1.7$, is detected in the full sample using the maximum likelihood method. The magnitude of evolution is at the lower end of previous estimates. The value of $P$ shows some dependence on the redshift boundaries which might be caused by different mixtures of infrared subpopulations which have very different evolutionary trends. The present-epoch 60 $\\mu$m luminosity function derived from galaxies in the redshift range $0.003<z<0.2$ shows good agreement with the luminosity function from Saunders et al. (1990). However, the luminosity function derived from galaxies in the redshift range $0.2<z<4.0$ shows evidence for much stronger evolution or possibly an entirely new class of infrared galaxies of high luminosity emerging at redshift $z>0.2$. Follow-up spectroscopy studies will investigate further the population at the high luminosity end.  The spatial clustering of galaxies as a function of galaxy properties such as type and luminosity provides strong constraints on models of galaxy formation and evolution. In particular, the clustering of the infrared galaxies contains essential information on the relation between star formation activity and environment. We use the projected correlation function $\\Xi(\\sigma)$, which is free from redshift-space distortion, to study the clustering dependence on infrared template type and infrared luminosity. When comparing clustering strength between the two infrared subpopulations, we found that cirrus-type galaxies cluster more strongly than M82 starbursts in all luminosity bins until the cirrus-population becomes absent at high luminosity/redshift. The average relative bias between the two infrared populations is $b_{\\textrm{cirrus}}/b_{\\textrm{M82}}=1.25\\pm0.07$ on scales $2\\le r \\le13~h_{72}^{-1}$ Mpc. This is consistent with the SFR - density relation which states that star-forming galaxies preferentially reside in low density environments. The dependence of clustering strength on SFR (using infrared luminosity as a proxy) is investigated using M82-type starbursts. We detect a correlation between SFR and clustering strength, which follows the basic trend found in the distant Universe (at $z\\sim1$). It can probably be explained by the tight correlation between SFR and stellar mass which holds at the local Universe as well as $z\\sim1$ and possibly $z\\sim2$. However, the trend that the SFR increases with clustering strength must becomes invalid at some point, as we do not find galaxies with the highest SFR in environments of highest density (such as centres of rich clusters) in the local Universe. The two seemingly inconsistent statements, the SFR - density relation and the correlation between SFR and clustering strength, can be reconciled if the latter operates mainly in the low density regime and becomes invalid once the density exceeds a certain threshold.  Direct measurement of the infrared emission of galaxies is crucial to measuring SFR of distant galaxies. We hope to extend our study of the relation between SFR and environment to higher redshift using surveys such as SWIRE and COSMOS. Observations to be made by the Herschel satellite will undoubtedly provide more insight."
    },
    "0910/0910.2970_arXiv.txt": {
        "abstract": "{Diffuse radio emission, in the form of radio halos and relics, traces regions in clusters with shocks or turbulence, probably produced by cluster mergers. The shocks and turbulence are important for the total energetics and detailed temperature distribution within the intracluster medium (ICM).  Only a small fraction of clusters exhibit diffuse radio emission, whereas a large majority of well-studied clusters shows clear substructure in the ICM. Some models of diffuse radio emission in clusters indicate that virtually all clusters should contain diffuse radio sources with a steep spectrum.  External accretion shocks associated with filamentary structures of galaxies could also accelerate electrons to relativistic energies and hence produce diffuse synchrotron emitting regions. The detection of radio emission from such filaments is important for our understanding of the origin of the Warm-Hot Intergalactic Medium (WHIM), and relativistic electrons and magnetic fields in the cosmic web. Here we report on Giant Metrewave Radio Telescope (GMRT) observations of a sample of steep spectrum sources from the 74~MHz~VLSS survey.    These sources are diffuse on scales~$\\gtrsim~15\\arcsec$, and not clearly associated with nearby ($z~\\lesssim~0.1$) galaxies. } { The main aim of the observations is to search for diffuse radio emission associated with galaxy clusters or the cosmic web. } {We have carried out GMRT~610~MHz continuum observations of unidentified diffuse steep spectrum sources. } { We have constructed a sample of diffuse steep spectrum sources, selected from the 74~MHz VLSS survey. We identified eight diffuse radio sources probably all located in clusters. We found five radio relics, one cluster with a giant radio halo and a radio relic, and one radio mini-halo. The giant radio halo has the highest radio power ($P_{1.4}$) known to date. By complementing our observations with measurements from the literature we find correlations between the physical size of relics and the spectral index, in the sense that smaller relics have steeper spectra. Furthermore, larger relics are mostly located in the outskirts of clusters while smaller relics are located closer to the cluster center. } { } ",
        "introduction": "Studies of large-scale structure (LSS) formation have made significant advances during the last decade. It has been found that nearly all massive clusters have undergone at least several mergers in their history and that presently clusters are still in the process of accreting matter. A significant fraction of the accreting mass is in the form of (smaller) clusters and galaxy groups. Cluster mergers are the most energetic events in the present day Universe, with kinetic energies of the order of $10^{63}-10^{64}$ erg, which are dissipated in giant shock waves and turbulence. An important aspect is the total energy budget and the detailed temperature distribution within the ICM, both of them are affected by the merger history of a cluster \\citep[e.g.,][]{2008SSRv..134..311D}.  Diffuse steep spectrum radio emission is observed in about 50 massive merging and post-merging galaxy clusters \\citep[see the review by][and references therein]{2008SSRv..134...93F}. This diffuse emission is difficult to detect due to its low surface brightness and steep spectral index\\footnote{$F_{\\nu} \\propto \\nu^{\\alpha}$, with $\\alpha$ the spectral index}, so the fraction of clusters hosting diffuse radio emission is probably larger than we currently know. The diffuse emission in clusters is commonly divided into three main classes \\citep{1996IAUS..175..333F}. \\emph{Radio Halos} are extended ($\\gtrsim 1$~Mpc) diffuse unpolarized ($\\lesssim 5\\%$) sources, located in the center of clusters. They have a regular smooth appearance,  and follow the thermal X-ray emission. \\emph{Radio Relics} are elongated structures with an irregular morphology, mostly located in the periphery of clusters. Relics can be highly polarized ($10-50\\%$). Several different subclasses have been identified \\citep{2004rcfg.proc..335K}. Most known radio relics and halos are found in clusters which show signs of a current or recent merger. This supports the scenario in which the relativistic electrons are accelerated by merger-induced shocks or turbulence. However, \\emph{Radio Mini-halos} are not associated with merging clusters. They are found in the centers of cool core clusters \\citep[e.g.,][]{1991A&ARv...2..191F,  2006PhR...427....1P} and are associated with the central cluster galaxy and typically have sizes $\\lesssim 500$~kpc, with the diffuse emission surrounding the central cluster galaxy \\citep[e.g.,][]{2009A&A...499..371G}.   Two different mechanisms for in-situ acceleration of particles have been proposed to explain relics in clusters: (i) adiabatic compression of fossil radio plasma by a passing shock wave producing a so called radio ``phoenix'' \\citep{2001A&A...366...26E, 2002MNRAS.331.1011E}, or (ii) diffusive shock acceleration (DSA) by the Fermi-I process \\cite[e.g.,][]{1983RPPh...46..973D, 1987PhR...154....1B, 1991SSRv...58..259J, 1998A&A...332..395E, 2001RPPh...64..429M}. In the first scenario, radio relics should have toroidal and complex filamentary morphologies. These relics are capable of producing very steep, curved radio spectra due to  inverse Compton (IC) and synchrotron losses. In the DSA scenario the electrons are accelerated by multiple crossings of the shock front (in a first order Fermi process). These relics have large sizes (Mpc) and are direct tracers of shock fronts in clusters. The spectral index is determined by the balance between the continuous acceleration at the shock front and energy losses in the post-shock region.  The  diffuse emission within clusters reveals the presence of relativistic electrons and magnetic fields on scales $\\sim 1$~Mpc. Spectral aging, due to synchrotron and IC losses of the emitting electrons may explain the steep spectra. If the electrons are injected via Fermi acceleration (DSA), their energy follows a power-law distribution. The power-law index of the injected electrons is related to the Mach number of the shocks \\citep[e.g.,][]{2007MNRAS.375...77H}: shocks with a low Mach number have steeper radio spectra. % Clearly, low-frequency surveys are needed to locate and study these sources \\citep{2006MNRAS.369.1577C, 2007MNRAS.378.1565C, 2008A&A...480..687C}. Interestingly, \\cite{2008Natur.455..944B} discovered a radio halo in the cluster \\object{Abell~521}, which was previously known to host a radio relic, with a spectral index of $\\sim-2.1$, suggesting the existence of a population of diffuse source in clusters with spectral indices~$< -1.5$.   Numerical simulations show the development of various types of shocks during structure formation~\\citep{2000ApJ...542..608M}. These shocks differ in their location with respect to the cluster center and Mach numbers \\citep{2000ApJ...542..608M, 2002MNRAS.337..199M, 2003ApJ...593..599R, 2006MNRAS.367..113P, 2009MNRAS.395.1333V}. External accretion shocks have $\\mathcal{M} \\gg 1$ and process the low-density, unshocked intergalactic medium (IGM). This results in relatively flat spectral indices of about $-0.5$ at the location of the shock front. Further away from the shock front the spectral index steepens due to synchrotron and IC losses. Internal shocks, (i.e., merger and flow shocks) occur within the cluster. The Mach numbers of these shocks are lower resulting in steeper spectral indices. Binary merger shocks are the result of a cluster merging with a another cluster or a large sub-structure. In this case a double radio relic is expected \\citep[e.g.,][]{1999ApJ...518..603R, 2008MNRAS.391.1511H, 2009arXiv0908.0728V}.    As pointed out by \\cite{1999ApJ...514....1C}, hydrodynamic models indicate that up to half of the baryons at present time should have temperatures in the range of $10^{5}-10^{7}$~K. Unfortunately, studying the abundance and distribution of this WHIM is very challenging, since its main tracers are highly excited Oxygen lines which are difficult to observe \\citep[e.g.,][]{2005Natur.433..495N}. A fraction of the accretion shocks will be supersonic and can accelerate energetic electrons up to energies of $10^{18}-10^{19}$ eV \\citep[e.g.,][]{1995ApJ...454...60N, 1996ApJ...456..422K, 2008ICRC....4..555I}. In the presence of magnetic fields, such electrons will emit faint diffuse synchrotron radiation.  The detection of these radio filaments is very important as this would provide a probe of the WHIM. Recent magnetohydrodynamical modeling indicates that detecting radio emission from the filamentary cosmic web should be possible \\citep[e.g.,][]{2004ApJ...617..281K, 2007MNRAS.375...77H}. \\cite{2008MNRAS.391.1511H, 2008MNRAS.385.1242P, 2008MNRAS.385.1211P} however find that in the outskirts of clusters (at a few times the virial radius) or filaments, external accretion shocks cause little radio emission, owing to the low density of both magnetic field energy and cosmic ray (CR) particles there \\citep{2001ApJ...562..233M}. They are therefore difficult to detect even with the sensitivity of upcoming radio telescopes such as LOFAR.  Relic emission from internal accretion shocks occur in a higher density environment so that they should be detected with current radio facilities.    When searching for radio halos, relics and filaments in low-frequency radio surveys, various other steep spectrum sources are also present. These include ultra-steep (angular size $\\lesssim 15\\arcsec$) spectrum sources \\citep[USS, see ][for a review]{2008A&ARv..15...67M} associated with high-z radio galaxies (HzRG), ``fossil'' or ``dying'' FR-I \\citep{1974MNRAS.167P..31F} radio sources, and ``head-tail'' sources, the last two having a steep spectrum due to spectral aging of the radio emission. In high-resolution ($\\lesssim 5\\arcsec$) observations, for example at $1.4$~GHz with the Very Large Array (VLA), diffuse objects will be resolved out due to missing short baselines. This provides a method for selecting diffuse radio sources associated with galaxy clusters or the cosmic web. FR-I sources can be partly excluded by removing sources that are clearly associated with individual galaxies.   The $74$~MHz VLA low-frequency Sky Survey (VLSS), \\cite{2007AJ....134.1245C}, covers about $3\\pi$~steradians of sky north of $\\delta=-30\\degr$. The resolution of the survey is 80\\arcsec~(FWHM) and the rms noise level is about 0.1~Jy~beam$^{-1}$. The source catalog contains roughly 70,\\ 000 sources with a point source detection limit of $0.7$~Jy beam$^{-1}$. A new calibration algorithm \\citep{2004SPIE.5489..180C} was used to remove the ionospheric distortions, which can be severe at this low-frequency. The $1.4$ GHz~NRAO VLA Sky Survey (NVSS), \\cite{1998AJ....115.1693C}, covers the entire sky above $\\delta = -40\\degr$. The NVSS images have a rms noise of about 0.45~mJy~beam$^{-1}$, and a resolution of 45\\arcsec (FWHM). The catalog contains about $2 \\times 10^{6}$ sources above a flux of $\\sim 2.5$~mJy.  In this paper we present radio continuum observations of 26~diffuse (angular size $\\gtrsim 15\\arcsec$) steep spectrum sources  selected from the VLSS survey with the GMRT at 610~MHz. The main aim of this project is to determine the morphology of the sources and search for diffuse structures which could be associated with shock fronts or turbulence in clusters, and accretion shocks onto filaments of galaxies.  The layout of this paper is as follows. In Sect.~\\ref{sec:selection} we discuss the sample selection, this is followed by an overview of the observations and data reduction in Sect.~\\ref{sec:obs-reduction}. In Sect.~\\ref{sec:results} we present the radio maps of the most interesting sources and discuss these sources individually. By combining our radio observations with data from the literature (X-ray and optical observations) we have tried to classify the sources. In Sect. \\ref{sec:spectra}, spectral indices are modeled using our flux measurements combined with literature values. We end with a discussion and conclusions in Sects.~\\ref{sec:discussion} and~\\ref{sec:conclusion}.  Throughout this paper we assume a $\\Lambda$CDM cosmology with $H_{0} = 71$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{m} = 0.3$, and $\\Omega_{\\Lambda} = 0.7$. ",
        "conclusions": "\\label{sec:conclusion} We carried out 610~MHz radio continuum observations with the GMRT of 26 diffuse steep spectrum sources. All of these sources were resolved out by 1.4~GHz~VLA~B-array snapshot observations or in the 1.4~GHz~FIRST survey. Of these 26 sources 25 were detected with the GMRT. % The radio observations show a wide range of sources: radio relics, a giant radio halo, a mini-halo, FR-I, head-tail, and USS sources. Here we shortly summarize the properties of the relics and halos in our sample.  -The radio relic 24P73 (VLSS~J2217.5+5943) is located close to the galactic plane. This source was previously known to have a steep spectrum but remained unclassified due to the limited sensitivity and resolution of previous observations. Due to the crowed Milky Way field and high extinction no cluster is found in POSS-II images. Deep NIR imaging will be needed to reveal the presence of a cluster. The sources has an extremely steep curved spectrum, likely the result of adiabatic compression of fossil radio plasma and spectral aging.  -VLSS~J1133.7+2324 is a filamentary relic located next to a nearby spiral galaxy. The radio source breaks up into two parallel filamentary structures. This relic source is probably associated with a structure of galaxies located at a distance of $z \\sim 0.6$. The nearby spiral galaxy itself is also detected in the radio images, the flux density consistent with that predicated by the FIR-radio correlation (using IRAS fluxes for this galaxy). The galaxy is partly blocking our view of the distant structure of galaxies, making follow-up observations more difficult.  -VLSS~J1431.8+1331 consists of two components. The eastern part is probably associated with a cD galaxy in the center of the cluster \\object{MaxBCG~J217.95869+13.53470} at $z=0.16$. The radio morphology suggests we are seeing signs of interaction between the radio plasma and the ICM. The other part is probably a relic source. These sources makes an interesting target for follow-up X-ray observations.  -VLSS ~J0004.9$-$3457 is associated with an elliptical galaxy in a small cluster at $z\\sim0.3$. The diffuse radio emission surrounds the galaxy in the form of a mini-halo. The source has an arc-like extension to the east.  - VLSS~J0717.5+3745 is located in the cluster MACS~J0717.5+3745, a massive merging system at a redshift of 0.5548. The radio emission is complex, consisting of a large elongated radio relic and a giant 1.2~Mpc radio halo. The halo has a radio power ($P_{1.4}$) of $5\\times 10^{25}$~W~Hz$^{-1}$, the highest one known to date. The relic might trace a giant shock wave related to the cluster merger. This is consistent with the very hot ICM and relatively flat spectral index.  -VLSS~J1515.1+0424 is a radio relic consisting of three filamentary structures in the periphery of the cluster Abell~2048.  Adiabatic compression of fossil radio plasma probably resulted in the complex morphology.\\\\  We find that larger relics are mostly located in the cluster periphery, while smaller relics are found closer to the cluster center. We also discovered a correlation between the spectral index and the physical size of the relics. A likely explanation for this correlation is that the larger shock waves occur mainly in lower-density regions and have larger Mach numbers. As a consequence they have shallower spectra. However, the correlation can also (partly) be explained by invoking different origins for relics. Relics formed by diffusive shock acceleration are direct tracers of large shock fronts in clusters, while the compression of fossil radio plasma is thought to produce relatively small relics with steep curved spectra.   If the correlation is explained by the fact that the larger shock waves occur mainly in lower-density regions and have larger Mach numbers then the injection spectral indices should flatten away from the cluster center for the various relics. To measure this requires a sufficient number of flux measurements, especially below the break frequency in the spectrum. With upcoming new radio facilities operating at lower frequencies it should therefore be possible to break the degeneracy between the injection spectral index and spectral aging.     Follow-up observations to measure the polarization properties are currently underway. We also plan to create spectral index maps for our sources which should give more information on the life-times of the various synchrotron emitting regions. Optical observations have been taken to characterize the large-scale environment around the radio sources. For two of our diffuse sources these observations should confirm the presence of a cluster. X-ray observations will be needed to study the dynamical state of the clusters and reveal the presence of shock fronts in the ICM."
    },
    "0910/0910.2141_arXiv.txt": {
        "abstract": "We study large--scale kinematic dynamo action due to turbulence in the presence of a linear shear flow, in the low conductivity limit. Our treatment is non perturbative in the shear strength and makes systematic use of both the shearing coordinate transformation and the Galilean invariance of the linear shear flow. The velocity fluctuations are assumed to have low magnetic Reynolds number ($\\rem$) but could have arbitrary fluid Reynolds number. The equation for the magnetic fluctuations is expanded perturbatively in the small quantity, $\\rem$. Our principal results are as follows: (i) The magnetic fluctuations are determined to lowest order in $\\rem$ by explicit calculation of the resistive Green's function for the linear shear flow; (ii) The mean electromotive force is then calculated and an integro--differential equation is derived for the time evolution of the mean magnetic field. In this equation, velocity fluctuations contribute to two different kinds of terms, the ``C'' and ``D'' terms, in which first and second spatial derivatives of the mean magnetic field, respectively, appear inside the spacetime integrals; (iii) The contribution of the ``D'' terms is such that their contribution to the time evolution of the cross--shear components of the mean field do not depend on any other components excepting themselves. Therefore, to lowest order in $\\rem$ but to all orders in the shear strength, the ``D'' terms cannot give rise to  a shear--current assisted dynamo effect;  (iv) Casting the integro--differential equation in Fourier space, we show that the normal modes of the theory are a set of shearing waves, labelled by their sheared wavevectors; (v) The integral kernels are expressed in terms of the velocity spectrum tensor, which is the fundamental dynamical quantity that needs to be specified to complete the integro--differential equation description of the time evolution of the mean magnetic field; (vi) The ``C'' terms couple different components of the mean magnetic field, so they can, in principle, give rise to a shear--current type effect. We discuss the application to a slowly varying magnetic field, where it can be shown that forced non helical velocity dynamics at low fluid Reynolds number does not result in a shear--current assisted dynamo effect. ",
        "introduction": "Large--scale magnetic fields in many astrophysical systems, from planets to clusters of galaxies, are thought to originate from dynamo action in the electrically conducting fluids in these objects. The standard paradigm involves amplification of seed magnetic fields, due to non mirror--symmetric (i.e. helical) turbulent flows, through the $\\alpha$--effect \\citep{Mof78, Par79}. It is only relatively recently that the role of the mean shear in the turbulent flows is beginnng to be appreciated. Dynamo action due to shear and turbulence has received some attention in the astrophysical contexts of accretion disks \\citep{VB97} and galactic disks \\citep{Bla98}. It has also been demonstrated that shear, in conjunction with rotating turbulent convection, can drive a large--scale dynamo \\citep{KKB08, HP09}. We are interested in the more specific problem of large--scale dynamo action due to  non--helical turbulence with mean shear.  Direct numerical simulations now provide strong support for such a {\\it shear dynamo}. \\citet{You08a} demonstrated that forced small-scale non--helical turbulence in non--rotating linear shear flows leads to exponential growth of large--scale magnetic fields. These findings were later generalized by \\citet{You08b} to a shearing sheet model of a differentially rotating disk with a Keplerian rotation profile. The investigations of \\citet{BRRK08} demonstrated the shear dynamo effect for a range of values of the Reynolds numbers and the shear parameter, and measured the tensorial magnetic diffusivity tensor. While the shear dynamo has been conclusively demonstrated to function, it is not yet clear what makes it work. This outstanding, unsolved problem is the focus of the present investigation.  One possibility that has been suggested is dynamo action due to a ``fluctuating $\\alpha$--effect'' in turbulent flows which have zero mean helicities. In this proposal, large--scale dynamo action derives from the interaction of mean shear with fluctuations of helicity \\citep{VB97,Sok97,Pro07,BRRK08,RK08,Sch08}. Another suggestion is that, if even transient growth makes non axisymmetric mean magnetic fields strong enough, they themselves might drive motions which could lead to subcritical dynamo action \\citep{ROPC08}. Yet another possibility that has been suggested is the shear--current effect \\citep{RK03,RK04,RK08}. In this mechanism, it is thought that the mean shear gives rise to anisotropic turbulence, which causes an extra component of the mean electromotive force (EMF), leading to the generation of the cross--shear component of the mean magnetic field from the component parallel to the shear flow. However, there is no agreement yet whether the sign of such a coupling is favourable to the operation of a dynamo. Some analytic calculations \\citep{RS06,RK06} and numerical experiments \\citep{BRRK08} find that the sign of the shear--current term is unfavourable for dynamo action. A quasilinear theory of dynamo action in a linear shear flow of an incompressible fluid which has random velocity fluctuations was presented in \\cite{SS09a,SS09b}. Unlike earlier analytic work which treated shear as a small perturbation, this work did not place any restriction on the strength of the shear. They arrived at an integro--differential equation for the evolution of the mean magnetic field and argued that the shear--current assisted dynamo is essentially absent. It should be noted that the quasilinear theory of \\cite{SS09a,SS09b} assumes zero resistivity, and is valid in the limit of small velocity correlation times when the  ``first order smoothing approximation'' (FOSA) holds.  In this paper we present a kinematic theory of the shear dynamo that is non perturbative in the shear strength, but perturbative in the magnetic Reynolds number ($\\rem$); this may be thought of as FOSA with finite resistivity. Thus we are not limited to the quasilinear limit of small velocity correlation times, and our conclusions are rigorously valid for velocity fluctuations which have small $\\rem$ but arbitrary fluid Reynolds number. In \\textsection~2 we formulate the shear dynamo problem for small $\\rem$. Using Reynolds averaging, we split the magnetic field into mean and fluctuating components. The equation for the fluctuations is expanded perturbatively in the small parameter, $\\rem$. Using the shearing coordinate transformation, we make an explicit calculation of the resistive Green's function for the linear shear flow. In \\textsection~3, the magnetic fluctuations and the mean electromotive force (EMF) are determined to lowest order in $\\rem$.  The transport coefficients are given in general form in terms of the two--point correlators of the velocity fluctuations. Galilean invariance is a basic symmetry in the problem and is the focus of \\textsection~4. For Galilean invariant (G--invariant) velocity fluctuations, it is proved that the transport coefficients, although space-dependent, possess  the property of translational invariance in sheared coordinate space. An explicit expression for the Galilean--invariant mean EMF is derived. We put together all the results in \\textsection~5 by deriving the integro--differential equation governing the time evolution of the mean magnetic field. Some important properties of this equation are discussed. In particular, it is shown that, in the formal limit of zero resistivity, the quasilinear results of \\cite{SS09a,SS09b} are recovered.  We also show that the natural setting for the integro--differential equation governing mean--field evolution is in sheared Fourier space. We prove a theorem on the form of the two--point velocity correlator in Fourier space; the velocity spectrum tensor and its general properties are discussed. We then express all the integral kernels in terms of the velocity spectrum tensor, which is the fundamental dynamical quantity that needs to be specified. Summary and conclusions are presented in \\textsection~6. ",
        "conclusions": "We have formulated the problem of large--scale kinematic dynamo action due to turbulence in the presence of a linear shear flow, in the limit of small magnetic Reynolds number ($\\rem$) but arbitrary fluid Reynolds number. The mean--field theory we present is non perturbative in the shear parameter, and makes systematic use of the shearing coordinate transformation and the Galilean invariance of the linear shear flow. Using Reynolds averaging, we split the magnetic field into mean and fluctuating components. The mean magnetic field is driven by the Curl of the mean EMF, which in turn must be determined in terms of the statistics of the velocity fluctuations. In order to do this it is necessary to determine the magnetic fluctuations in terms of the mean magnetic field and the velocity fluctuations. So we develop the equation for the fluctuations perturbatively in the small parameter, $\\rem$. Using the shearing coordinate transformation, we make an explicit calculation of the resistive Green's function for the linear shear flow. From the perturbative scheme it is clear that the fluctuations can be determined to any order in $\\rem$. Here we determine the magnetic fluctuations and the mean EMF to lowest order in $\\rem$. The transport coefficients are given in general form in terms of the two--point correlators of the velocity fluctuations. At this point we make use of Galilean invariance, which is a fundamental symmetry of the problem. For Galilean invariant velocity statistics we prove that the transport coefficients, although space-dependent, possess the property of translational invariance in sheared coordinate space. An explicit expression for the Galilean--invariant mean EMF is derived.  We put together all the results in \\textsection~5 by deriving the integro--differential equation governing the time evolution of the mean magnetic field. Some important properties of this equation are the following:  \\begin{enumerate} \\item Velocity fluctuations contribute to two different kinds of terms, the ``C'' and ``D'' terms, in which first and second spatial derivatives of the mean magnetic field, respectively, appear inside the spacetime integrals.  \\item The ``C'' terms are a generalization to the case of shear, of the ``$\\alpha$'' term familiar from mean--field electrodynamics in the absence of shear. However, they can also contribute to ``magnetic diffusion''; see discussion below. Likewise, the ``D'' terms are a generalization to the case of shear, of the ``magnetic diffusion''  term familiar from mean--field electrodynamics in the absence of shear. It must be noted that the generalization is non perturbative in the shear strength.  \\item In the mean--field induction equation, the ``D'' terms are of such a form that: (i) the equations for $B_1$ or $B_3$ involve only $B_1$ or $B_3$, respectively; (ii) the equation for $B_2$ involves both $B_1$ and $B_2$. Therefore, to lowest order in $\\rem$ but to all orders in the shear strength, the ``D'' terms cannot give rise to  a shear--current assisted dynamo effect.  \\item In the formal limit of zero resistivity, the quasilinear theory of \\cite{SS09a, SS09b} is recovered. In this case, the ``C'' terms vanish when the velocity field is non helical. However, this may not be the case when the resitivity is non zero. Whether the ``C'' terms give rise to such a shear--current--type effect depends on the form of the velocity correlators, which will be strongly affected by shear and highly anisotropic; hence it is difficult to guess their tensorial forms {\\it a priori} and it is necessary to develop a dynamical theory of velocity correlators -- see below for further discussion.  \\item Sheared Fourier space is the natural setting for the mean magnetic field;  the normal modes of the theory are a set of shearing waves, labelled by their sheared wavevectors.  \\item We prove a theorem (in the Appendix) on the form of the two--point velocity correlator in Fourier space; the velocity spectrum tensor and its general properties are discussed. The integral kernels are expressed in terms of the velocity spectrum tensor, which is the fundamental dynamical quantity that needs to be specified to complete the integro--differential equation description of the time evolution of the mean magnetic field. \\end{enumerate}  The physical meaning of the  ``C'' and ``D'' terms becomes clear in the limit of a slowly varying magnetic field, when the integro--differential equation reduces to a partial differential equation \\citep{SS10}. Then we  encounter the well--known $\\alpha$--effect and turbulent magnetic diffusion ($\\eta$), albeit in tensorial form. The ``C'' terms alone contributes to $\\alpha$, whereas both ``C'' and ``D'' terms contribute to magnetic diffusion. When the velocity field is non helical, the velocity spectrum tensor is real, and the tensorial $\\alpha$ coefficient vanishes; this result is true for arbitrary values of the shear parameter. The ``C'' terms can, in principle, contribute to a shear--current effect, through the off--diagonal components of the diffusivity tensor (which couple the streamwise component of the mean magnetic field with the cross--stream components). It turns out that these off--diagonal components depend on the microscopic resistivity in such a manner that they vanish when the microscopic resistivity vanishes. This result is consistent with the results of \\cite{SS09a, SS09b}. To deal with the case when the microscopic resistivity does not vanish, it is necessary to provide our kinematic development with a dynamical model for the velocity field. \\cite{SS10} show that, for forced non helical driving at low fluid Reynolds number, the sign of the off--diagonal terms of the diffusivity tensor does not favour the shear--current effect. This conclusion agrees with those reported in \\cite{RS06,RK06,BRRK08}, even if our results are limited to low Reynolds numbers. If we seek a different explanation for the dynamo action seen in numerical simulations, the  ``fluctuating $\\alpha$--effect'' still remains a promising candidate. $\\alpha$ itself is described by second--order velocity correlators, so to describe fluctuatons of $\\alpha$, it is necessary to deal with either fourth--order velocity correlators or products of two second--order velocity correlators. This requires extending our perturbative calculations by at least two higher orders, a task which, while tractable,  is beyond the scope of the present investigation."
    },
    "0910/0910.0691_arXiv.txt": {
        "abstract": "Nonlinear evolution of circularly polarized Alfv\\'en waves are discussed by using the recently developed Vlasov-MHD code, which is a generalized Landau-fluid model. The numerical results indicate that as far as the nonlinearity in the system is not so large, the Vlasov-MHD model can validly solve time evolution of the Alfv\\'enic turbulence both in the linear and nonlinear stages. The present Vlasov-MHD model is proper to discuss the solar coronal heating and solar wind acceleration by Alfve\\'n waves propagating from the photosphere. ",
        "introduction": "Large-amplitude Alfv\\'enic fluctuations are ubiquitous in the heliosphere, especially in the fast solar wind (Bruno and Carbone, 2005). Nonlinear evolution and dissipation of these fluctuations are thought to play important roles in heating of the solar wind plasmas (Suzuki and Inutsuka, 2006; Wu and Yoon,2007; Araneda \\etal, 2008; Valentini \\etal, 2008) and generation of the localized structures (Tsurutani \\etal, 2005; Vasquez \\etal, 2007; Lin \\etal, 2009). Since these fluctuations are typically robust for linear ion-cyclotron damping due to their small wave frequencies and for linear Landau damping due to their small propagation angle relative to the background magnetic field, wave-wave interactions (parametric instabilities) are significant processes of their nonlinear evolution.  The quasi-parallel propagating Alfv\\'enic fluctuations can resonate with both the parallel and obliquely propagating magnetohydrodynamic (MHD) waves (Mjolhus and Hada, 1990; Champouex \\etal, 1999; Nariyuki \\etal, 2008). In uniform plasmas with a typical parameter of the solar wind, parallel propagating Alfv\\'en waves dominantly resonate with the parallel propagating Alfv\\'en waves and ion acoustic waves. It is important that Alfv\\'enic fluctuations can resonate with the \\lq\\lq kinetic'' wave modes such as ion acoustic waves. Namely, even if the fluid approximation is valid for Alfv\\'enic fluctuations themselves, the ion kinetics should be considered for parametric instabilities (Araneda, 1998; Bugnon \\etal, 2004; Nariyuki and Hada, 2006; 2007, Araneda \\etal, 2007).  Furthermore, in the plasmas with inhomoginity and/or unstable velocity distribution functions such as beam components, the resonance and dissipation processes can be different from those in the homogenous plasmas (Tsikrauri \\etal, 2005; Suzuki and Inutsuka, 2005; 2006; Wang \\etal, 2006; Suzuki, 2008; Nariyuki \\etal, 2009). The kinetic Alfven waves which propagate at oblique angles relative to the background magnetic field can be excited by the proton beams in the solar wind (Daughton and Gary, 1998; Yin \\etal, 2007). We note that such beam-excited waves can preferentially interact with the finite-amplitude parallel propagating Alfv\\'en waves. Actually, the observational studies imply the importance of the kinetic Alfv\\'en waves on the dissipation of the Alfv\\'enic turbulence (Leamon \\etal, 1998; Hamilton \\etal, 2008).  As mentioned above, the nonlinear evolution and dissipation processes of Alfv\\'enic fluctuations in the solar wind are \\textit{cross scale coupling processes}, in which the MHD-, ion-, and electron-scale phenomena coexist and mutually interact. Moreover, it is worth noting that the heating and acceleration process of plasmas by Alfv\\'enic fluctuations are desired to be slot into the heliospheric global simulation models (\\eg, Nakamizo \\etal, 2009). Thus, the development of the alternative \\lq\\lq kinetic'' MHD models, which are the \\lq\\lq middel scale'' model including some non-MHD effects, is necessary for \\textit{systematic} understanding on such cross scale coupling processes.  As one of such models, we have recently developed a new Vlasov simulation code named Vlasov-MHD code (1-D in the configuration space, 1-D in \\lq\\lq kinetic '' velocity space, and 2-D in \\lq\\lq MHD '' velocity space), in order to study basic properties of nonlinear evolution of Alfv\\'en waves (Kumashiro \\etal, 2009). The concept of the Vlasov- MHD model is the generalized model of the so-called Landau fluid model (Passot and Sulem, 2003; Bugnon \\etal, 2004), which includes the nonlinear wave-particle interactions. The linear analysis of the Vlasov-MHD model has been carried out (Araneda, 1998; Nariyuki and Hada, 2006; 2007, Araneda \\etal, 2007) and has concluded that as far as the amplitude of the parent waves is not so large, the growth rates of the linear analysis are consistent with those in the numerical results of the ion hybrid simulation. Kumashiro \\etal (2009) performed the numerical simulation using the Vlasov-MHD code with the Hall-effect (Vlasov-Hall-MHD code) and demonstrated that the linear growth of parametric instabilities of Alfv\\'en waves are almost consistent with those of the ion hybrid simulation.  In the present paper, we discuss the nonlinear evolution of Alfv\\'en waves using the Vlasov-Hall-MHD code. In section 2, we briefly introduce the numerical schemes of Vlasov-Hall-MHD code. We present the simulation results in section 3. Section 4 summarizes the results and briefly discuss the future issues. ",
        "conclusions": "In the present paper, we discussed the nonlinear evolution of finite amplitude Alfv\\'en waves using the Vlasov-MHD simulation code, which is a generalized Landau fluid model including the nonlinear wave-particle interactions. It was confirmed that while the Landau damping should be evaluated along perturbed field lines (Finn and Gerwin, 1996), as far as the nonlinearlity of parent Alfv\\'en waves are weak, numerical results of the present Vlasov-MHD simulation agree well with those of the ion hybrid simulations both in the linear and nonlinear stage. It is worth noting that while the computational load of Vlasov-MHD code is much less than that of ion hybrid simulation. On the other hand, when the nonlinearlity of parent Alfv\\'en waves are not weak, the present Vlasov-MHD simulation is not proper. Furthermore, the Vlasov-MHD model is also inadequate for the high ion beta plasmas, since the finite Larmor radius effects are neglected.  We note that in spite of restriction on nonlinearlity and ion kinetics, the present Vlasov-MHD model can be applied to several heliospheric problems. One of them is the solar coronal heating / solar wind acceleration by Alfv\\'en waves propagating from the photosphere (Suzuki and Inutsuka, 2006). Since the amplitude of the magnetic fluctuations is not so large and the beta ratio is small (Suzuki and Inutsuka, 2006; Tanaka \\etal, 2007), the assumption in the Vlasov-MHD system is well satisfied.  The characteristics of the observed velocity distribution function near the sun (\\eg Marsch, 2006) is one of the important constraint on the solar coronal heating and solar wind acceleration model. Namely, the generation processes of the heliospheric nonequilibrium plasmas such as ion beam components, temperature anisotropy, and the relative speed among the ion spices should be comprehensively examined in the model. We note that such a point of view is important for the collaborating works among future missions (BepiColombo, SCOPE/Cross-Scale, Solar-C, Solar Orbiter, and so on), which achievements will contribute toward the heliospheric science."
    },
    "0910/0910.5282_arXiv.txt": {
        "abstract": "We present the results of a search for occultation events by objects at distances between 100 and 1000~AU in lightcurves from the Taiwanese-American Occultation Survey (TAOS). We searched for consecutive, shallow flux reductions in the stellar lightcurves obtained by our survey between 7 February 2005 and 31 December 2006 with a total of $\\sim4.5\\times10^{9}$ three-telescope simultaneous photometric measurements. No events were detected, allowing us to set upper limits on the number density as a function of size and distance of objects in Sedna-like orbits, using simple models. ",
        "introduction": "During and just after the epoch of planet formation, a significant number of planetesimals were left that had not accreted into planets. Some of the planetesimals near the giant planets were ejected to orbits with large semi-major axes, and many more were hurled hyperbolically into interstellar space. Those objects which formed beyond proto-Neptune would probably remain near the ecliptic plane and became part of the Kuiper Belt \\citep{1950BAN....11...91O, 1974CeMec...9..321K, 1980Icar...42..406F, 1987AJ.....94.1330D}.   Up to the present, more than 1000~Kuiper Belt Objects (KBOs) have been discovered\\footnote{See \\url{http://www.cfa.harvard.edu/iau/lists/TNOs.html} for a list of these objects.}. The estimated total mass of the Kuiper Belt from observations to date is $\\sim 0.1~$M$_{\\oplus}$, which is a few orders of magnitude less than the mass predicted for the minimum mass solar nebula \\citep{1977Ap&SS..51..153W}.  Theoretical simulations show that in order to assemble 1000-km size objects between 30 and 50~AU, initially there must have been much more mass in this region \\citep{2004Natur.432..598K, 2005AJ....129..526S} than the current observations suggest. Many surveys indicate a Kuiper Belt ``cliff'' beyond 50~AU \\citep{1999AJ....118.1411C, 2001AAS...199.6310A, 2004AJ....128.1364B}. Given that more than 99\\% of the original disk mass appears to have been depleted, some sort of dynamical perturbations must have removed objects from this region.  Furthermore, the origin of $2003~$VB$_{12}$ (Sedna, with orbital parameters of $a$=531~AU, $q$=76~AU, $i=12^\\circ$) is puzzling \\citep{2004ApJ...617..645B}. Searches for thermal radiation from Sedna with the IRAM 30~m telescope and Spitzer have put an upper limit on its diameter of $1600$ km \\citep{2004DPS....36.0301B}. The mass of Sedna, from the upper limit on the size and mean density of KBOs, is estimated to be $\\leq 10^{-3}~$M$_{\\oplus}$. Given its large size and highly eccentric orbit, in situ formation seems unlikely.  Sedna is unique because its perihelion at 76~AU is far outside the reach of dynamical perturbation by Neptune, and the aphelion of $\\sim$1000~AU is also too small for Sedna to be affected sufficiently by giant molecular clouds or galactic tides in the solar neighborhood. Meanwhile, recent observations \\citep{2005A&A...439L...1B,2007A&A...466..395E} have shown that Sedna is likely to have methane, water and nitrogen ices.  The similarity of the spectroscopic features between Sedna and Triton \\citep{2005A&A...439L...1B} suggests its formation near the giant planet region.  Several theories have been proposed to explain the existence of Sedna in its present orbit. For instance, a close stellar encounter \\citep{2000Icar..145..580F, 2000ApJ...528..351I, 2004Natur.432..598K, 2004AJ....128.2564M} could have depleted the Kuiper Belt and excited outer Solar System objects, such as Sedna, to their present orbits. Other possibilities have also been suggested, including rogue mars-sized bodies, or multiple stellar encounters in the Sun's birth cluster \\citep{2004ApJ...617..645B, 2004AJ....128.2564M, 2006ApJ...643L.135G, 2007Icar..191..413B}.  Note that all of these theories also predict that many more objects would be found in \\emph{Sedna-like} orbits, i.e. orbits with perihelia large enough to be unaffected by Neptune and aphelia small enough to be unaffected by galactic tidal forces or giant molecular clouds. Furthermore, the fact that Sedna was discovered very close to perihelion, where it spends a very small fraction of its orbital period, led \\citet {2004ApJ...617..645B} to infer that there may be as much as many as 500 objects in similar orbits with the same size as Sedna or larger. However, recent search attempts \\citep{2007AJ....133.1247L, 2008ssbn.book..335B, schwamb2009} did not find any new member of this population. In their recent survey, \\citet{schwamb2009} modeled objects with the same semi-major axis and eccentricity as Sedna.  A best fit of 40 objects brighter than or equal to Sedna in that region is consistent with their results of detecting one object with a perihelion past 70~AU.  Most objects in Sedna-like orbits would have escaped from direct detection due to their enormous distances from the Earth.  Direct detection of this kind of object is limited to relatively close objects and by large telescopes, so the majority of the population is still well beyond the detection limit for most ground-based telescopes. \\citet{1976Natur.259..290B} suggested that the existence of the ``invisible'' objects in this region might be inferred by stellar occultation, and the utility of using robotic telescopes to search for occultations by outer Solar System objects has been demonstrated \\citep{axelrod1992, Cook1995}.   The Taiwanese-American Occultation Survey (TAOS) aims to investigate the size distribution of small ($\\sim$1~km) KBOs \\citep{2009PASP..121..138L, 2008ApJ...685L.157Z}.  Four 50~cm robotic telescopes have been installed at Lulin mountain in central Taiwan.  Each telescope is a fast F/2 system with a 3 deg$^{2}$ field of view, equipped with a 2k$\\times$2k CCD camera that can perform 5~Hz photometric sampling on a few hundred stars simultaneously. Multiple telescopes are used to reduce the false positive event rate.  The TAOS system has been in routine operation with three telescopes since February 2005, and four-telescope observations began in August 2008.  TAOS was designed to search for KBOs near 50~AU, for which a typical occultation by a KBO of a few km across would manifest itself as a flux reduction in only one or two consecutive time series measurements.  With a null detection, a stringent constraint has been set on the number and size distribution of KBOs by \\citet{2008ApJ...685L.157Z}.  However, the same data set can be used to detect objects in the distant Solar System all the way to $\\sim$1000~AU; i.e., the survey is also sensitive to objects in Sedna-like orbits.  For objects at 100 -- 1000~AU, the Fresnel scale and the projected size of the background star are larger than for objects in the Kuiper Belt, resulting in occultation events with longer durations \\citep{2007AJ....134.1596N}. Thus a distant occultation event would have flux reductions spanning \\emph{several} lightcurve points and requires a different analysis pipeline on the TAOS data from that reported in \\citet{2008ApJ...685L.157Z}.  Here we present our analysis of the TAOS data to search for objects on Sedna-like orbits.  An estimate of the number of such objects detected would help constrain the population of bodies beyond 50~AU, thereby confronting theoretical models of the early Solar System, and address the puzzle of mass deficit in the region.  \\begin{figure*}[b] \\plottwo{f1a.eps}{f1b.eps} \\caption{Diffraction profiles for 10~km diameter object occulting a 0.01 mas background star (projected stellar diameter $\\sim7$~km at 1000~AU) at 200~AU and 1000~AU. Solid lines are the diffraction profiles, solid circles represent 0.2~sec (5~Hz) sampling light curve points, dotted lines indicate the duration of the 0.2 sec integration. The sampling starts at random time offset. The velocity of the object was calculated at an opposition angle of $30^{\\circ}$.} \\label{fig:diffraction} \\end{figure*} ",
        "conclusions": "We have analyzed the first two years (2005 February to 2006 December) of TAOS data to search for objects in Sedna-like orbits. With the null detection of objects in this region, we are able to set upper limits upon the number density of objects at different distances. Given that one Sedna has been found, we have also placed lower limits on the number density of objects at 100~AU. We also show that the candidate events reported in \\citet{2006AJ....132..819R} are inconsistent with the TAOS data at the 95\\% c.l. if the size distribution at $\\Delta=200$ AU has slope $q < 7.6$. We plan to modify the detection algorithm to facilitate an exact calculation of the statistical significance of candidate events, and we will subsequently complete the analysis of the second TAOS data set (January 2007 to December 2008) and report its results in a future paper. We will also expand the analysis to include limits on the populations of Scattered Disk and Inner Oort Cloud objects."
    },
    "0910/0910.0002_arXiv.txt": {
        "abstract": "The coalescence of supermassive black hole binaries occurs via the emission of gravitational waves, that can impart a substantial recoil to the merged black hole.  We consider the energy dissipation, that results if the recoiling black hole is surrounded by a thin circumbinary disc.  Our results differ significantly from those of previous investigations.  We show analytically that the dominant source of energy is often potential energy, released as gas in the outer disc attempts to circularize at smaller radii.  Thus, dimensional estimates, that include only the kinetic energy gained by the disc gas, underestimate the real energy loss.  This underestimate can exceed an order of magnitude, if the recoil is directed close to the disc plane.  We use three dimensional Smooth Particle Hydrodynamics (SPH) simulations and two dimensional finite difference simulations to verify our analytic estimates.  We also compute the bolometric light curve, which is found to vary strongly depending upon the kick angle.  A prompt emission signature due to this mechanism may be observable for low mass $(10^6 \\ M_\\odot)$ black holes whose recoil velocities exceed $\\sim 10^3 \\ {\\rm km \\ s}^{-1}$. Emission at earlier times can mainly result from the response of the disc to the loss of mass, as the black holes merge. We derive analytically the condition for this to happen. ",
        "introduction": "Supermassive black hole binaries are predicted by hierarchical galaxy formation models and are a likely consequence of observed galaxy mergers.  However, only a handful of these binaries have been directly observed \\citep{rodriguez06} and their dynamical evolution is still uncertain.  If binaries coalesce on a time-scale shorter than the age of the universe, mergers can be an important ingredient in the evolution and growth of supermassive black holes. Mergers also emit low frequency gravitational waves whose detection is one of the prime goals of proposed experiments such as the {\\em Laser Interferometer Space Antenna} (LISA). What remains uncertain is whether detectable electromagnetic emission occurs either prior to or in the immediate aftermath of the coalescence.  Observations and numerical simulations strongly suggest that mergers of gas-rich galaxies result in the inflow of gas into the nuclear region, where it is likely to co-exist with a newly formed black hole binary \\citep{escala05,dotti07,callegari09}. The gas may have an active role and mediate and speed up the coalescence \\citep{ivanov99,natarajan02,cuadra09,lodato09}. However -- and irrespective of whether the gas is {\\em dynamically} important -- if gas is present then perturbations to the gas during the merger may give rise to observable  electromagnetic signals that can either precede or follow the gravitational wave signal.  Observations of these signals can give us evidence for the existence of mergers, improve upon LISA's limited ability to spatially localise sources, and teach us about the astrophysical phenomena involved in the process.  General relativity predicts that gravitational waves emitted by the shrinking binary during the final coalescence carry away a non-zero net linear momentum, so that the centre of mass of the merged black holes recoils \\citep{peres62,baker07}. A number of authors have used semi-analytic methods or N-body simulations to estimate the influence of the recoil on a surrounding thin disc \\citep{lippai08,shields08,schnittman08}. Hydrodynamic simulations for a thick disc -- a configuration that we do not consider here -- have also been performed by \\cite{megevand09}.  Here we revisit the analytic problem (emphasizing the importance of different physical effects from those discussed previously), and perform numerical simulations to investigate how a thin, non self-gravitating disc reacts to the birth of a central black hole via merger. The initial disc geometry that we consider is based upon that discussed by \\cite{natarajan02} and \\cite{milos05}. We assume that, at a time significantly prior to the final coalescence (when the binary semi-major axis was $> 10^2 $ Schwarzschild radii), the binary was surrounded by a geometrically thin disc whose plane coincided with that of the binary. While the binary remains in tidal contact with this disc, energy and angular momentum can be exchanged between the two. This exchange provides an additional source of heat for the inner disc \\citep{lodato09}, and alters its long term viscous evolution \\citep{pringle91}. Immediately prior to the coalescence, however, the rapid decay of the binary orbit due to gravitational wave emission leads to a decoupling of the binary from the inner disc. As a result, it is a reasonable first approximation to model the effect of a kick on an initially unperturbed, axisymmetric circumbinary disc. We also ignore any effects associated with gas that has survived in circumprimary or circumsecondary discs up until near the moment of coalescence. Such gas, if present, may produce observable signatures \\citep{natarajan02, lodato09,chang09}, but it is decoupled dynamically from the gas in the circumbinary disc. The effect of the recoil on the circumbinary disc can therefore be considered independently from the fate of gas bound to either of the individual black holes.   In this paper, our main focus is on the effect of the black hole recoil rather than the almost instantaneous mass loss that also accompanies the merger \\citep{milos05}.  We study the consequence of velocity perturbations in both the linear (kick velocity smaller than circular velocity) and non linear regimes.  We calculate the magnitude of energy dissipation and assess the dependence of this potentially observable quantity on the recoil geometry. In particular, we focus on off-plane kicks, which are expected to be the most common.   The paper is organised as follows. In Section~\\ref{sec:ana} we analytically investigate the properties of the disc after the kick and the dissipation of the extra energy in the innermost region. These estimates serve as guidance for our numerical simulations, described in Section~\\ref{sec:sim}. In Section~\\ref{sec:results} we present our numerical results, which we discuss in Section~\\ref{sec:discussion}. Finally, in Section \\ref{sec:conclusion} we draw our conclusions. ",
        "conclusions": "\\label{sec:conclusion} In this paper we have investigated the effect of a post-merger kick on the dynamics of a geometrically thin accretion disc, surrounding the newly merged black hole. The existence and strength of kicks following black hole mergers are now secure predictions of General Relativity. Whether small-scale gas discs typically attend mergers remains, however, uncertain. If gas discs are commonly present the perturbation to the gas caused by the kick will generate an electromagnetic counterpart to the merger, which may be detectable if its amplitude and time scale are observable at cosmological distances.  The main conclusions of our study are: \\begin{itemize} \\item[1.] The dimensional assumption that the magnitude of energy release is $(1/2) V^2 \\Sigma_{\\rm V} R^2_{\\rm V}$  is an underestimate. The energy available for dissipation varies strongly with the angle between the kick and the disc plane, and for kicks close to the plane our analytic estimate exceeds $(1/2) V^2 \\Sigma_{\\rm V} R^2_{\\rm V}$ by up to three orders of magnitude. A large increase in the energy release for kicks close to the plane is confirmed numerically. \\item[2.] For most orientations of the kick, accretion energy (i.e. energy liberated when gas in the disc loses angular momentum and falls inward) dominates over the direct energy input to the gas in the frame of the kicked black hole. Accretion energy is particularly important for kicks that are directed toward the disc plane. The importance of accretion can be demonstrated analytically, and is confirmed by numerical simulations. \\item[3.] We have run SPH and (for the special cases that are two-dimensional) finite difference numerical simulations to investigate the evolution of the disc and the rate of energy dissipation, following a kick at an arbitrary angle to the disc plane. The simulations yield explicit predictions for the form of the light curves as a function of both kick angle and surface density structure within the disc. \\item[4.] The observability of emission from kicked discs depends upon the kick velocity, the orientation of the kick relative to the disc plane, and the mass of the disc at the radii where the energy is deposited. We have argued that the mass of the disc is limited by the requirement that the disc remains stable to fragmentation due to self-gravity. If this is true, the decreased disc mass (as compared to that assumed in prior works) outweighs the effect of the increased energy release per unit mass. As a result, the predicted luminosity around massive black holes is substantially smaller than previously thought. It is unlikely that such sources can be identified via wide-area sky surveys. \\item[5.] The most feasible observational probe of the phenomena discussed here is via identification of variable disc emission following black hole mergers detected by other means (e.g. via detection of gravitational waves). Disc emission may be detectable provided that the kick velocity is large, especially in the case where the kick is directed close to the disc plane. To quantify, a kick velocity of $10^3 \\ {\\rm km \\ s}^{-1}$ offers promising possibilities, but $300 \\ {\\rm km \\ s}^{-1}$ appears to be very difficult or impossible to detect.  \\end{itemize} Physically, the magnitude and orientation of the kick is determined by the magnitude and direction of the spin of the black holes immediately prior to their merger. Our results emphasize that the observability of this class of electromagnetic counterparts depends sensitively on the distribution of the pre-merger spins. Studies, such as that by \\citet{bogdanovic07} and that by \\citet{king08}, that attempt to predict the evolution of black hole spin during the earlier phases of merger are thus particularly important for determining whether counterparts will be detectable. Determining the spectral signature of the kicked disc is also important, and this will require simulations of the disc evolution that include more complete treatments of the disc physics."
    },
    "0910/0910.3001_arXiv.txt": {
        "abstract": "\\baselineskip=15pt  We present the two-dimensional (2D) ionization structure of self-similar magnetohydrodynamic (MHD) winds off accretion disks around irradiated by a central X-ray point source. Based on earlier observational clues and theoretical arguments, we focus our attention on a subset of these winds, namely those with radial density dependence $n(r)\\propto 1/r$ ($r$ is the spherical radial coordinate). We employ the photoionization code \\verb\"XSTAR\" to compute the ionic abundances of a large number of ions of different elements and then compile their line-of-sight (LOS) absorption columns. We focus our attention on the distribution of the column density of the various ions as a function of the ionization parameter $\\xi$ (or equivalently $r$) and the angle $\\theta$. Particular attention is paid to the absorption measure distribution (AMD), namely their Hydrogen-equivalent column per logarithmic $\\xi$ interval, $d N_H/d \\log \\xi$, which provides a measure of the winds' radial density profiles. For the chosen density profile $n(r)\\propto 1/r$ the AMD is found to be independent of $\\xi$, in good agreement with its behavior inferred from the X-ray spectra of several active galactic nuclei (AGNs). For the specific wind structure and X-ray spectrum we also compute detailed absorption line profiles for a number of ions to obtain their LOS velocities, $v \\sim 100 - 300$ km~s$^{-1}$ (at $\\log \\xi \\sim 2-3$) for \\fexvii~ and  $v \\sim 1,000 - 4,000$ km~s$^{-1}$ (at $\\log \\xi \\sim 4-5$) for \\fexxv, in good agreement with the observation. Our models describe the X-ray absorption properties of these winds with only {\\it two parameters}, namely the mass-accretion rate $\\dot m$ and LOS angle $\\theta$. The probability of obscuration of the X-ray ionizing source in these winds decreases with increasing $\\dot m$ and increases steeply with the LOS inclination angle $\\theta$. As such, we concur with previous authors that these wind configurations, viewed globally, incorporate all the requisite properties of the parsec scale ``torii\" invoked in AGN unification schemes. We indicate that a combination of the AMD and absorption line profile observations can uniquely determine these model parameters and their bearing on AGN population demographics. ",
        "introduction": "The issue of accretion as the power source behind Active Galactic Nuclei (AGNs) was decided as early as 1969 \\citep[][]{LyndenB69}. Since then, a 40-year effort to unravel the physics underlying this process has produced a great body of observational and theoretical work covering a host of the AGN properties and phenomena. However, due to the small angular size of a black hole horizon, whether at the galactic or extragalactic setting, these studies were effectively performed mainly in the spectral (but also the time) domain; the spatial structure of accretion flows, but those of scale of many parsecs, was then inevitably delegated to models, whose validity was determined by their ability to reproduce the spectral and timing observations with sufficient accuracy. At the risk of oversimplifying the issue, the spectral properties of AGNs were basically outlined (among other works) in \\citet{Sanders89} who showed that the Spectral Energy Distribution (SED) of AGNs includes three broad components in the IR, optical-UV and X-ray bands, with roughly equal energy per decade, with the radio comprising only a small fraction of their bolometric luminosity, of order $\\sim 10^{-6}$ in the Radio-Quiet (RQ) AGNs and $\\sim 10^{-3}$ in Radio-Loud (RL) AGNs.  Considering that accretion produces most of its power in the last 10 or so Schwarzschild radii, $r_S$, outside the black hole horizon, barring an inherently non-thermal emission process (as it turns out to be the case with blazars), it is strange that the observed SEDs exhibit roughly constant luminosity per decade. Interestingly, the simplest of assumptions, namely radiation of the AGN bolometric luminosity in black body form by an object (most likely in the form of a cold disk of $T \\sim 10^4 - 10^6$ K) of size a few $r_S$ of a supermassive black hole ($M \\sim 10^6-10^8 \\, M_{\\odot}$ where $\\Msun$ is the solar mass), implies peak emission at UV frequencies, consistent with the observationally identified UV feature referred to above, known as the Big Blue Bump (BBB); as such, this  feature was modeled with emission by a geometrically-thin, optically-thick accretion disk extending usually to the Innermost Stable Circular Orbit (ISCO) of the black hole. Of the other two distinct components the X-rays are generally attributed to emission by an optically-thin, hot ($T \\sim 10^8-10^9$ K) corona overlying this disk and heated by the action of magnetic fields that thread the geometrically-thin disk, in close analogy to the Solar Corona \\citep[e.g.][]{Haardt91,Kawanaka08}. Finally, the IR emission was accounted for as the result of reprocessing the O-UV radiation by a geometrically thick (scale height $h$ roughly equal to the local radius $r$, i.e. $h \\sim r$), cool, molecular torus at large distances ($\\gsim 1$ pc) from the central engine [an arrangement that presents a bit of a problem, since the temperature of this torus ($\\sim 10 - 100$ K) is much less than the virial temperature of the gas at that distance and would lead to a configuration with $h \\ll r$; see Krolik \\& Begelman~(1988). The region between the torus and the accretion disk was thought to be occupied by clouds with velocity widths ranging from tens of thousands to several hundreds of kilometers per second, producing the observed line radiation, with the ensemble of these components encapsulated in the well known cartoon of \\citet[][]{UP95}.  While in most models these components are generally thought of as independent, Principal Component Analysis (PCA) of multiwavelength AGN features by \\cite{BorGre92} and \\cite{Bor02}  indicated that most of the variance in the measured optical emission line properties and a broad range of continuum features (radio, optical, X-ray) of AGNs was contained in two sets of correlations, eigenvectors of the correlation matrix. This analysis suggested that the properties of these features are not independent (despite their diverse locations and emission processes) but are well correlated by two single underlying physical parameters through the physics of accretion and the conversion of its power to radiation and outflows. On the other hand, general theoretical arguments suggesting, e.g. a tight correlation between variations in the X-ray and O-UV components, given their implied proximity, were not confirmed by observations \\citep{Nandra98}.    The advent of X-ray instrumentation and in particular X-ray spectroscopy, established the presence of spectral components apparently present across the entire range of compact object masses from galactic black hole candidates (GBHCs) to luminous AGNs, namely outflowing X-ray absorbing matter in our line of sight (LOS) manifested as complex (blueshifted) absorption features of hydrogen equivalent column density $N_H \\sim 10^{21} - 10^{23}$ cm$^{-2}$. The presence of this absorbing material, often referred to as Warm Absorber, was first established in {\\it ASCA} observations of many AGNs \\citep[e.g.][]{ReyFab95} and GBHCs \\citep[e.g.][]{Miller06} which provided clear evidence of highly ionized oxygen \\ovii ~(0.739 keV) and/or \\oviii ~(0.871 keV) along the LOS. More recently, grating spectra of high-resolving power obtained by {\\it Chandra} and {\\it XMM-Newton} have enabled the study of these absorption features in greater detail leading to the conclusion that they are present in a large fraction of AGNs and span a wide range ($\\sim 10^5$) in ionization parameter \\citep[e.g.][hereafter HBK07]{HBK07}. Some sources appear to contain low and/or high velocity outflows \\citep[e.g. HBK07;][]{HBA09} while a handful of objects exhibit a trans-relativistic outflow \\citep[e.g.][]{Chartas09,Pounds09,Reeves09}. In addition in both AGNs and GBHCs the observed outflows from at least a number of sources have been suggested to be accelerated magnetically (rather than by another process) off an accretion disk (e.g., see Miller et al.~2006, 2008 for GRO~J1655-40; Kraemer et al.~2005 and Crenshaw \\& Kraemer~2007 for NGC~4151). A brief review of luminous AGN winds is found in \\citet{Brandt09}. These facts imply the presence of gas covering a large fraction of the solid angle and distributed over a large range of radii, perhaps the entire range between the X-ray source and the torus, indicating AGNs to be multiscale and multiwavelength objects rather than a class with properties determined only by the power released by accretion in the black hole vicinity. This last point has gained further support with the discovery of the correlation between the accreting black hole mass and the velocity dispersion of the surrounding stellar population \\citep{FerMer}.   From the theoretical point of view, most early AGN treatments were focused on the structure of the innermost regions of the accretion flow and as such they limited themselves to the study of the conditions in a rather narrow range of radii. Nonetheless, several treatments of magnetohydrodynamic (MHD) accretion disk winds, aiming to account for the observed AGN outflows, employed self-similar solutions, which naturally span a large range in radius (e.g. Blandford \\& Payne 1982, hereafter BP82; Contopoulos \\& Lovelace 1994, hereafter CL94). Using the structure of these solutions, \\citet[][hereafter KK94]{KK94} proposed that the so-called molecular torus of the AGN unification scheme is a dynamical rather than a static object, namely an MHD wind of the type suggested in the above works. KK94 also contented that agreement between model and observation demanded a rather specific type of wind, a particular case of those discussed in CL94, one that we also concentrate on in the present work. At the same time, they noted that the 2D geometry of these winds (with low column at inclination angles $\\theta \\simeq 0^{\\circ}$, i.e. along the wind axis, and high column at $\\theta \\simeq 90^{\\circ}$) provided, in addition, a natural framework for the unification scheme of type I and II Seyfert galaxies proposed by Antonucci \\& Miller~(1985), thereby linking the large scale AGN spectral classification to the dynamics of AGN accretion/outflows.  Self-similar solutions of the accretion flow equations were also provided by \\citet{NY94,NY95a}, who considered the structure of hot accretion flows in the regime of low accretion rate, i.e. for accretion rates less than the Eddington value, $\\dot m \\equiv \\dot M/ \\dot M_{\\rm E} \\lsim 1$. Because for these low rates the accretion time scale $\\tau_{acc}$ is shorter than the gas cooling time $\\tau_{\\rm cool}$, in fact $\\tau_{\\rm acc} \\simeq \\dot m \\, \\tau_{\\rm cool}$, only a fraction $\\dot m$ of the energy released in the dissipation of the plasma kinetic energy is radiated away with the remainder advected into the black hole (hence the term Advection Dominated Accretion Flows or ADAFs); as a result, the accretion luminosity is then proportional to $\\dot m^2$ (rather than $\\dot m$). The high temperature of the inner ADAFs, provides naturally for the hot component required to produce the observed X-ray emission, which, rather than being an independent component, within ADAFs is incorporated in the dynamics of accretion itself \\citep[see][for recent review]{Narayan08}. More recently, \\citet{BlandBegel}, elaborated further on ADAFs elucidating, among others, the reason for which the Bernoulli integral of ADAFs is positive \\citep[a point that had been noted by][]{NY94,NY95a}, that being that the viscous stresses transfer outward, in addition to angular momentum, also mechanical energy. They then proposed that this excess energy can be carried away in the form of wind over a large range of radii, thereby leading to configurations similar to those obtained in the more detailed studies of BP82 and CL94, while at the same time maintaining some of the general properties of ADAFs. The combination of radiatively inefficient flows with the simultaneous presence of winds led to the nomenclature Advection Dominated Inflow-Outflow Solutions (ADIOS). Following these pioneering works, many attempts have been made in recent years to reproduce and explain the kinematics and X-ray spectra of the observed outflows in AGNs \\citep[see, among others,][]{Proga03,Everett05,Sim05,SD07,DKP08,Sim08,SD09}.  Motivated by these considerations we are taking a closer look at the issue of Warm Absorbers (X-ray absorbing medium) in terms of specific disk-wind models, namely those of CL94. As it will be discussed in more detail later on, these solutions are characterized by a parameter $q$, which determines the distribution of axial current in the wind as a function of the radius; this parameter, in conjunction with the MHD conservation laws and the radial force balance, determines both the radial dependence of the matter density in the wind and the radial dependence of the toroidal magnetic field component $B_{\\phi}$. Taking our lead from the works of \\citet{Behar03} and HBK07, in the present work we restrict our attention to wind models with $q=1$, i.e. the value used also by KK94. This is in a sense a critical value of this parameter because, as discussed in CL94, it leads to winds with toroidal magnetic field $B_{\\phi} \\propto 1/r$ and therefore magnetic field energy on the disk that diverges only logarithmically at small and large $r$ and as such it can be considered as a ``minimum magnetic energy\" configuration (KK94); values of $q \\neq 1$ lead to configurations with power-law divergence either at small or large $R$. In addition, the choice $q=1$ leads to a wind with radial density profile $n(r) \\propto r^{-1}$, i.e. a profile with equal column per logarithmic radial interval (and normalization that depends on the polar angle $\\theta$ and the mass-accretion rate $\\dot m$), a fact that is in general agreement with the profiles implied by the ionized absorbers data analysis to date (HBK07). It is worth noting that a similar density profile was invoked by \\citet{KHT96} and \\citet{HKT97} to explain the Fourier frequency dependence of the soft-hard X-ray lags observed in GBHCs and in AGNs (Papadakis et al. 2002).   With the wind 2D ($r$ and $\\theta$) density field as a function of radius and angle provided by the models of CL94 (i.e. ignoring at present effects that can be attributed to other agents that could affect the wind structure) and the assumption of a point X-ray source of a given spectrum at the origin, one can compute their 2D ionization structure and in particular the local column density of specific ionic species as a function of radius and angle. One can also compute the integrated column of each such ion, quantities that can be directly compared to observation (HBK07). This we do in the present work. { While a number of issues remain open or are sidestepped in the present self-similar wind models, this is the price to pay in order to limit the number of free parameters to just two (the wind mass flux and the observer inclination angle), a fact that allows the study of the properties of warm absorbers within the global perspective of AGN unification}.   Our paper is structured as follows: in \\S 2 we review the MHD wind equations and structure as given in CL94 and make a connection of these winds with flows of the ADAF/ADIOS type by relating the wind density normalization to the ionizing luminosity produced by the idealized point source at the coordinate origin. In \\S 3 we provide the general 2D ionization structure, the local column of specific ions as a function of radius as well as the total ionic column of a given charge state and for each element, as well as the absorption line profiles of selected transitions. In \\S 4 the results are compared to observation, the properties and limitation of the present treatment are discussed and a course for future work is charted. ",
        "conclusions": "In this work we have presented a detailed study of the ionization structure of model MHD winds off accretion disks; this is our first attempt to model within this context the recent observations of absorption features in the {\\it ASCA, XMM} and {\\it Chandra} X-ray spectra, the so-called warm absorbers. To this end we have employed the self-similar 2D hydromagnetic wind models developed by CL94 that provide the fluid 2D density and velocity fields, which we coupled to photoionization calculations using \\verb\"XSTAR\". Our consideration of magnetocentrifugal (MHD) winds/outflows in this work has been motivated and supported in part by recent observational implications that at least in a number of AGNs and GBHCs the inferred driving mechanism of the observed X-ray ionized wind medium is magnetic rather than thermal or radiative (e.g., see Miller et al.~2006, 2008 for GRO~J1655-40; Kraemer et al.~2005 and Crenshaw \\& Kraemer~2007 for NGC~4151).  Given the scope of our paper, our model wind is necessarily overly simple. It ignores a host of issues that affect the structure of winds off accretion disks and replaces them with the self-similar models of CL94. The interested reader can get a feeling of the multitude and complexity of the issues not addressed in the present treatment by looking at the works of, e.g., \\citet[][]{Proga03} and \\citet[][]{ProgKal04} (and references therein) and also \\citet[][for disks with several distinct values of the accretion rate $\\dot {m}$]{Ohsuga09}. These include, among others, the radiative transfer in the wind, the ensuing effects of radiation pressure and also the influence of the vertical gradient of $B_{\\phi}^2$ in the launching of the wind. As shown in these references the winds can be radiation driven in the inner part of the disk (due to the increased flux at this spatial domain) while being magnetically launched at larger radii. In this respect our model winds are quite different in structure and variability from the more realistic winds of \\citet{Proga03} at small radii, while they should be more similar to those of his at larger radii where radiation pressure is smaller and the wind is magnetically launched. Realistic models must by necessity consider the transition of the underlying disk from radiation pressure, at small $r$, to gas pressure dominance at larger radii; this fact would likely force a different choice of our boundary condition $\\psi'(90^{\\circ})$ and therefore yield a field geometry that breaks the self-similarity of our solutions.  Wind models very similar to those used herein but with the inclusion of the effects of the radiation pressure that break the self-similarity were discussed by KK94 and \\citet{Everett05}.  One should therefore view our models with the above caveats in mind. The significance and assumptions of the models we propose is basically justified {\\em a posteriori} by their ability to interpret the observations. Self-similarity is one of the fundamental assumptions of our model winds. However, whether exactly self-similar or not, any model that would attempt to account the entire range of ionic species shown e.g. in \\citet{Behar09} must by necessity cover a very broad range in the photoionization parameter $\\xi$. Then, the column of each such ion provides a measure of the corresponding hydrogen equivalent column $N_H$ as a function of $\\xi$, i.e. the AMD. These two quantities ($N_H, \\xi$) can then be employed to provide a measure of the gas density $n(r)$ as a function of the distance $r$ from the X-ray source. Barring the possibility of several independent regions at different distances but similar columns as indicated by the functional form of the AMD (one could consider this is possibility at the risk of introducing an inordinate number of free parameters), models with radial density profiles as those considered here are a natural consequence of the AMD obtained by the {\\it Chandra} observations.      The wind ionization structure was followed in 1D as described in \\S 2.2, namely along the observer's LOS, assuming the ionizing source to be point-like. Clearly, a more comprehensive treatment of this problem should take into account also the backward emitted radiation which will reach the observer from the regions of the disk on the other side of the black hole, as well as the scattered radiation, whose effects could be significant. We plan to return to these issues in a future publication.  The 2D structure of the winds considered in our analysis implies, at a minimum, a two-parameter description of the absorption features. However, the self-similarity of the problem simplifies further the treatment by allowing the separation of the $r$ and $\\theta$ variables with the wind density obtaining the form  $n(r, \\theta) \\propto r^{2q-3} {\\cal N}(\\theta)$. The parameter $q$ is a most important parameter of these models because it determines both the radial dependence of the ionization parameter $\\xi$ and the wind column density per decade of radius along a given direction (LOS). In the models examined in the present work we restricted our study to the value $q=1$, which leads to a radial density profile $n(r) \\propto 1/r$. This is an interesting profile in that it provides for an ionization parameter $\\xi$ with a similar dependence, i.e. $\\xi \\propto 1/r$, and most importantly, with equal column per decade in radius, across the entire range in radius and the corresponding range in $\\xi$; as a result, ionic species of very different ionization properties, existing over a wide range of ionization parameter (and radius), have roughly comparable column densities, independent of the distance and are therefore possible to detect. As noted in CL94, the value of $q$ determines also the axial current distribution in radius within the wind, with the corresponding magnetic energy per unit length at the wind base being also constant per (cylindrical) decade in radius. It is worth comparing our work with that of KK94, who used the same type of MHD wind with the value of $q$ as we do herein: while these authors focused their study on the effects of dust on radiation transfer, we focus our attention on the effects of the ionizing radiation on the X-ray spectra of AGNs, in particular on the ionized absorbers. KK94 also noted in passing the relevance of their models to the X-ray AGN spectra well ahead of the detailed outflow observations made with {\\it Chandra}.   While the structure of the wind in $r$ and $\\theta$ is provided by the models of CL94, for the remaining parameters of our models, namely the normalization of the wind density and the X-ray luminosity of the ionizing source we have chosen to use the scalings of ADAFs, or better ADIOS; as a result the X-ray luminosity is proportional to the square of the (normalized) accretion rate at the inner edge of the disk $\\dot m^2$ rather than simply $\\dot m$, assumed in standard accretion disks. Therefore, restricting ourselves to models with $q=1$, the global ionization structure of a given wind, including the normalization of the column density, depends only on {\\it two parameters}, namely $\\dot m$ and the observer's inclination angle $\\theta$. This parametrization provides an extremely economical set of assumptions concerning not only the outflows, i.e. the wind ionization structure seen in absorption in detailed X-ray spectra, but also for the entire (radio-quiet) AGN unification picture: Figure \\ref{fig:density} which exhibits the normalization of the wind density as a function of $\\theta$ along the magnetic field line of $\\Psi = \\Psi_0$, makes apparent the difference in column between face-on and edge-on views, implying that the wind, if extending to sufficiently large radii, can in fact serve as the proposed molecular torus associated with AGN unification, a point originally made by KK94, also for winds with $q=1$; it is of interest to note that the few objects for which sufficiently detailed observations exit are consistent values $q \\simeq 1$ \\citep[HBK07;][]{Behar09}. Within this same context and ionizing luminosity considerations, one should note the dependence of source obscuration on the X-ray luminosity, namely its reduced value for objects accreting at a higher fraction of their Eddington rate $\\dot m$, which apparently is in general agreement with observations. It remains to be seen whether these notions can withstand the scrutiny of more consistent and encompassing observational tests.   At this point we would like to stress the importance of the AMD in the study of AGN outflows/winds, a quantity enunciated by HBK07 and modeled in detail in the preceding sections. Analyses similar to those of HBK07 and \\citet{Behar09} are indispensable because they produce a consistent analysis of the entire set of absorption features in the AGN X-ray absorption spectra. At the same time they underscore the unique value of X-ray spectroscopy which, in a wavelength band of $\\sim 1.5$ decades, can encapsulate the properties of ions that span $\\sim 5$ orders of magnitude in ionization parameter and, for models with $q \\simeq 1$, a similar range in radius; at the same time, measurement of their absorption columns yields the equivalent hydrogen column of the flow, $N_H$, over a similar range in radius, thereby going a long way toward the determination of the physics underlying the outflow dynamics. The fact that the AMD analyses to date are roughly independent of $\\xi$ provide support to our use of the $q=1$ or $n(r) \\propto 1/r$ models, reiterating that a similar density profile was invoked on the basis of AGN and GBHC timing properties \\citep{PNK01,KHT96}.  We have thus presented a concrete example of the AMD dependence on $\\xi$ for our $q=1$ models with $\\epsilon = 0.2, ~ \\dot m = 0.1$ and two different values of the observer's inclination angle (see table~\\ref{tab:tbl-1} for details). As expected we found the AMD to be constant, i.e. independent of $\\xi$, over many decades in this parameter (i.e. the local column density $\\Delta N_H$ is independent of ionization states $\\xi$ of ions). With $\\dot m = 0.1, ~\\eta_W \\simeq 0.5$, equation~(\\ref{eq:column1}) implies $\\Delta N_H \\sim 2.6 \\times 10^{21}$ cm$^{-2}$ (yielding a total column of $N_H \\sim 3.9 \\times 10^{22}$ cm$^{-2}$) for $\\theta=30\\degr$ and $\\Delta N_H \\sim 1.8 \\times 10^{22}$ cm$^{-2}$ (total of $N_H \\sim 2.5 \\times 10^{23}$ cm$^{-2}$) for $60\\degr$ over $-1 \\lesssim \\log \\xi \\lesssim 4$ erg~cm~s$^{-1}$, in good agreement with the observed AMD of IRAS~13349+2438 \\citep[HBK07;][]{Behar09} and also some other AGNs detailed below. One should note that these values depend primarily on $\\dot m$ and $\\theta$ and are independent of the black hole mass. The mass of the object gets involved only as a measure of its total luminosity, which does not appear in the expression for $N_H$, implying that these models could in principle be applicable also in GBHCs. High quality X-ray absorption data are in fact available for the GBHC GRO~J1655-40 and were used to argue for magnetic driving of the wind in this system too \\citep{Miller06,Miller08}. These spectra are distinguished from those of AGNs by the prominent absence of low ionization state ions. This is to be expected given that the presence of the companion star limits the extent of the disk to roughly half the distance between the two objects or $r \\simeq 10^{12}$ cm. Given that the Schwarzschild radius of the compact object is $r_S \\simeq 10^6$ cm the entire disk size covers only a range of $\\log x \\sim 6$ in radius; considering (based on Fig.~\\ref{fig:amd}) that in the inner region of the wind ($r \\sim 1000-3000 r_S \\sim 10^9$ cm) major elements (even heavier species) are almost fully ionized, one would expect the presence of ions over only a factor of 1000 in $\\xi$, in rough agreement with the observation.   Given that our models provide also the complete velocity field of the MHD wind we have also produced a sequence of synthetic absorption profiles as described in \\S 2.2 and shown in \\S 3.2. We have done so for two charge levels of Fe, namely \\fexxv~ and \\fexvii. With our fiducial model predicting the \\fexvii~ column to be maximum at $\\log \\xi \\sim 2.2-3$ erg~cm~s$^{-1}$, we then infer the corresponding LOS velocity (see Fig.~\\ref{fig:fe}) to be $v_{\\rm los} \\sim 100-300$ km~s$^{-1}$ in excellent agreement with the observed values of IRAS~13349+2438. It is important to note that the model produces not only the correct velocity $v_{\\rm los}$ at the maximum value of $\\Delta N_{\\rm Fe}(\\xi)$ but also the observed $\\Delta N_H(\\xi)$ normalization {\\em for the same inclination angle}. It is therefore possible with measurements of the combined line widths and column densities to reproduce within the present models both $q$, $\\theta$ and $\\dot m$, thereby providing a complete specification of these winds. In addition to IRAS~13349+2438, the nearby bright Seyfert MCG~6-30-15 has been observed to show, as one of the multiple ionization zones, an X-ray ionized absorber with an outflow velocity of $\\sim 1900$ km~s$^{-1}$ at $\\log \\xi \\sim 3.85$ and $N_H \\simeq 9 \\times 10^{22}$cm$^{-2}$ \\citep[e.g.][]{Young05, HBA09} also in good agreement with our model results for $\\dot m \\sim 0.1$ (see Fig.~\\ref{fig:fe}). Clearly, a different choice of a set of conserved quantities (wind variables) in the model would produce a slightly different field line geometry in which the resulting MHD outflow could in principle obtain higher LOS velocities (perhaps by factors of magnitudes), and this will be studied in detail in a future work.  As noted above our model is quite successful in reproducing the observed AMD of IRAS~13349+2438. More recently, \\citet{Behar09} presented a compilation of the AMD of a number of AGN with $d N_H/d \\log \\xi$ which are slightly different from constant but with $q-$values of $n(r) \\propto r^{2q-3}$ which are still very close to unity; e.g. NGC~3783 ($q \\sim 0.89$), NGC~5548 ($q \\sim 0.94$), NGC~7469 ($q \\sim 0.9$) and MCG-6-30-15 $(q \\sim 0.95)$. We believe some of these AMDs are sufficiently close to those of our fiducial model to be virtually indistinguishable. By comparison, the MHD wind models of BP82 have $q=0.75$ and $n(r) \\propto r^{-3/2}$. Their ionization parameter decreases more slowly with $r$ ($\\xi \\propto r^{-1/2}$) and hence a range of 5-6 decades in $\\xi$ implies a range of 10-12 decades in radius with the corresponding distances being unrealistic. Most importantly, the wind columns would decrease with radius ($N_H \\propto r^{-1/2}$) and the AMD dependence on $\\xi$ would be $d N_H/d \\log \\xi \\propto \\xi$, in clear disagreement with these observations, and thus their model is essentially ruled out \\citep{Behar09}.   Despite the apparent success of this first model, one should bear in mind that several aspects have been treated in a rather simplified fashion. Here we discuss some of them and their influence on the results presented so far:  1. The shape of the ionizing spectrum used so far ($F_{\\nu} \\propto \\nu^{-\\alpha}$ between 13.6 eV and 13.6 keV with $\\alpha = 1.5$) is quite simplistic, however, such spectra are often used in similar type calculations \\citep[e.g.][]{Sim08}. This is a crude approximation to the observed spectra characterized by complex spectral shapes. Spectral features like the BBB and the UV to X-ray luminosity ratio $\\alpha_{\\rm OX}$ \\citep[][]{Elvis00} play a role in our results \\citep[a recent review of the broad band AGN SEDs can be found, e.g., in][]{Risaliti04}. More recently, \\citet{Grupe04} sampled 110 bright soft-Xray selected AGNs for a simultaneous study of Optical/UV and X-ray data and reported strong correlations between the X-ray spectral slope and the Optical/UV slope which should be included in a more comprehensive treatment. Multi-components of the broad-band AGN spectra, as those discussed by \\citet{Elvis00} and \\citet{Risaliti04}, should play an important role in characterizing the observed X-ray spectral features. Adding softer photons (responsible for the Optical/UV components) could impact the radiative transfer between those more complex photon distributions and the ionized matter \\citep[see, e.g.,][for multi-component injected spectra]{Everett05,Sim05}. In fact, we have mimicked the presence of such softer photons by choosing larger values of the spectral index of our models ($\\alpha \\gtrsim 2$). For the same total luminosity (implemented in the definition of $\\xi$), it was found in this trial run that the radial position of the peak column density of a given ion (e.g. \\fexvii) decreases, because the presence of a given ion requires a certain flux of {\\em ionizing} photons per atom; in a steep spectrum the proper ratio is found at larger values of the parameter $\\xi$ or correspondingly smaller values of the distance $x=r/r_S$. As a result, the widths of the corresponding transitions should be larger and therefore the proper ionizing spectrum is necessary for the correct interpretation of the relations between ionic column densities and wind kinematics \\citep[see][for a similar approach]{Everett05,Sim05}.   2. The ionizing source has been considered, for simplicity, to be point-like. A source of finite size will provide for more complex illumination, especially for parts of the wind at which the source angular extent is significant. Most importantly an extended source would impact the absorption line profiles, since a line of sight may pass through matter other than that used in section 3.2. Finally, to treat the radiative transfer correctly at all energies, one should eventually need to resort to a 2D approach rather than the 1D model (radial only) presently used, which would modify the global ionization structure (Fig.~\\ref{fig:2d}b).   3. The wind equations, as presently implemented include only MHD forces, while it is apparent that in the presence of the ionic species we produce one should also include the effects of radiation pressure. Because the latter depends on the ionization state of the plasma which in turn depends on its kinematics, implemented correctly, this should be done iteratively in a way similar to \\citet{Everett05}. Based on the rather small effect that the radiation force seems to have on the structure of these winds, we believe that our present results are generally valid.  In summary, we have presented above a first attempt at interfacing theoretical models of MHD winds with AGN observations, in particular the absorption features in their X-ray spectra of ionized outflows. We have placed most of our emphasis on modeling the corresponding AMDs, which, based on the existing observations seems to favor a specific value of the model parameter $q$ which determines the radial density dependence of our models, namely $q \\simeq 1$. With the value of this parameter set by observation, our models present a {\\it two-parameter} family of AGN wind structure, namely $\\dot m$ and $\\theta$. We have found that, based on the limited number of objects discussed above, our models fare rather well in accommodating the observations and the possibility of incorporating AGN unification within the same models, as discussed in KK94. There is still a multitude of issues related to X-ray absorption that remain open, e.g. whether the model can accommodate the high velocity ($v/c \\sim 0.1 - 0.25$) outflows associated with X-ray absorption features \\citep[see, e.g.,][for some attempts with Monte Carlo simulations]{Sim05,Sim08} in the spectra of some bright quasars such as APM~08279+5255, PG~1211+143 and PDS~456 \\citep[e.g.][]{Chartas09,Pounds09,Reeves09}. In the context of our model this point may be well explained with an optimized magnetic field geometry [i.e. $\\Psi(r,\\theta)$ distribution determined by GS-equation~(\\ref{eq:GS})] that provides a favorable LOS velocity component. Important as they are, these go beyond the scope of the present work and we expect to return to them in future publications."
    },
    "0910/0910.1835_arXiv.txt": {
        "abstract": "{% In this study we provide the first numerical demonstration of the effects of turbulence on the mean Lorentz force and the resulting formation of large-scale magnetic structures. Using three-dimensional direct numerical simulations (DNS) of forced turbulence we show that an imposed mean magnetic field leads to a decrease of the turbulent hydromagnetic pressure and tension. This phenomenon is quantified by determining the relevant functions that relate the sum of the turbulent Reynolds and Maxwell stresses with the Maxwell stress of the mean magnetic field. Using such a parameterization, we show by means of two-dimensional and three-dimensional mean-field numerical modelling that an isentropic density stratified layer becomes unstable in the presence of a uniform imposed magnetic field. This large-scale instability results in the formation of loop-like magnetic structures which are concentrated at the top of the stratified layer. In three dimensions these structures resemble the appearance of bipolar magnetic regions in the Sun. The results of DNS and mean-field numerical modelling are in good agreement with theoretical predictions. We discuss our model in the context of a distributed solar dynamo where active regions and sunspots might be rather shallow phenomena.}  \\begin{document} ",
        "introduction": "Turbulence effects generally refer to the occurrence of corre\\-la\\-tions between velocity, temperature, and/or magnetic fields at small scales. A typical example is turbulent viscosity, which results from the spatial exchange of turbulent eddies characterized by velocity correlations. This leads to the dissipation of energy at small scales.  However, there is also the possibility of additional (e.g.\\ non-diffusive) turbulence effects, as is perhaps best known in mean-field electrodynamics and dynamo theory. Here one models the effects of the mean electromotive force, i.e.\\ the turbulence effects of velocity and magnetic field fluctuations, on the evolution of the mean field. This can lead to the occurrence of the famous $\\alpha$ effect, in addition to turbulent magnetic diffusion, turbulent diamagnetic velocity, etc. (Moffatt 1978; Krause \\& R\\\"adler 1980). Another example is the $\\Lambda$ effect in rotating anisotropic hydrodynamic turbulence, which can lead to the occurrence of differential rotation in cosmic bodies such as the Sun (R\\\"udiger 1980, 1989; R\\\"udiger \\& Hollerbach 2004). In that case the relevant correlations come from the mean Reynolds stress tensor and its dependence on the local angular velocity.  A related example is the combined Reynolds and Max\\-well turbulent stress tensor and its dependence on the mean magnetic field. The first analytic calculations of the dependence of the turbulent Reynolds stress on the mean magnetic field in the framework of the first-order smoothing approximation were performed by R\\\"adler (1974) and R\\\"udiger (1976). Later, also the combined effects of the turbulent Reynolds and Maxwell stress tensors were considered (Kleeorin et al.\\ 1989, 1990; R\\\"udiger \\& Kichatinov 1990). It was noticed that this can lead to a local reduction of the total turbulent pressure and hence to the possibility of self-induced concentrations of large-scale magnetic fields (Kleeorin et al.\\ 1989, 1990, 1996; Kleeorin \\& Rogachevskii 1994; Rogachevskii \\& Kleeorin 2007). Such a process may play an important role in the formation of sunspots and active regions in the Sun. It may be complementary to the magnetically induced suppression of the turbulent energy flux, which would lead to further cooling and hence a further concentration of the structures (Kitchatinov \\& Mazur 2000).  The current leading explanation for the formation of sun\\-spots is related to the emergence of deeply rooted magnetic flux tubes (Parker 1955, 1982, 1984). Such flux tubes are generally believed to be produced and `stored' near the bottom of the convection zone (Spiegel \\& Weiss 1980). The storage of magnetic fields and the formation of flux tubes in the overshoot layer near the bottom of the solar convective zone was investigated in a number of publications (see, e.g., Spruit 1981; Spruit \\& van Ballegooijen 1982; Sch\\\"{u}ssler et al.\\ 1994; Moreno-Insertis et al.\\ 1996; Tobias et al.\\ 2001; Tobias \\& Hughes 2004). However, in order that the tubes retain their basic east--west orientation during their ascent over many pressure scale heights, the magnetic field must be strong enough (Choudhuri \\& D`Silva 1990) and is estimated to be around $10^5\\G$ at the bottom of the convection zone (D`Silva \\& Choudhuri 1993). Such fields would be up to a hundred times stronger than the equipartition value, which is one of several other arguments that have led to the idea that flux emergence of dynamo-generated fields might instead be a shallow phenomenon (Brandenburg 2005; Schatten 2009). Such a scenario appears also compatible with solar subsurface flows, as inferred from local helioseismology (Zhao et al.\\ 2001; Kosovichev 2002). In particular, Zhao et al.\\ (2004) and Hindman et al.\\ (2009) find the presence of converging flows around active regions at radii as large as 100--200\\Mm. It appears that these convergent flows might actually be the source of the formation of active regions and perhaps sunspots rather than a consequence (Parker 1979a; Hurlburt \\& Rucklidge 2000). Of course, in the immediate proximity of individual spots one observes outgoing flows. Those are probably superficial, less than 1--2\\Mm\\ deep, and appear to be due to the dynamical effects of magnetoconvection in an inclined magnetic field of the penumbra (e.g., Thomas et al.\\ 2002; Heinemann et al.\\ 2007; Rempel et al.\\ 2009; Kitiashvili et al.\\ 2009).  The goal of this paper is to investigate the effects of turbulence on the mean Lorentz force by means of direct numerical simulations (DNS) for forced turbulence and to study the instability of a uniform large-scale magnetic field in an adiabatically stratified layer by means of mean-field numerical modelling based on parameterizations both of analytic formulae by Rogachevskii \\& Kleeorin (2007) and the results of DNS for forced turbulence. In order to study the essence of the effect, we make several simplifications by neglecting the energy equation, i.e.\\ the specific entropy is assumed to be strictly constant in space and time. In the mean-field numerical modelling we neglect the suppression of turbulent magnetic diffusivity and turbulent viscosity, and omit correlations with density fluctuations. Nevertheless, the mean density is allowed to evolve fully self-consistently according to the usual continuity equation. ",
        "conclusions": "In this study we have demonstrated in DNS the effects of turbulence on the mean Lorentz force. This effect is quantified by determining the relevant functions $q_{\\rm p}(\\meanB)$ and $q_{\\rm s}(\\meanB)$ that relate the sum of the turbulent Reynolds and Maxwell stresses with the Maxwell stress of the mean magnetic field. Using three-dimensional simulations of forced hydromagnetic turbulence with an imposed field, we confirm that the function $q_{\\rm p}(\\meanB)$ is positive and can reach values much larger than unity for $\\meanB/\\Beq\\ll1$. This thereby reverses the sign of the effective magnetic pressure $P_{\\rm eff}(\\meanBB)$ associated with the mean magnetic field, which then becomes $P_{\\rm eff}(\\meanBB)= \\half(1-q_{\\rm p})\\meanB^2$. We find that the function $q_{\\rm s}(\\meanBB)$ that determines the modification of magnetic tension, also reaches values much larger than unity, but its value is less than half the value of $q_{\\rm p}$ and the error bars are larger. This work has also demonstrated explicitly the possibility of a large-scale instability of the full system of mean-field MHD equations. Finally, our study has revealed the presence of spatial structures arising from the instability that might be associated with the formation of bipolar magnetic regions and perhaps also sunspots when applied to the Sun.  The effects of turbulence on the large-scale Lorentz force may also be important in applications to the solar torsional oscillations by changing the mean magnetic tension $\\meanBB \\cdot \\nab \\meanBB$ (R\\\"{u}diger et al.\\ 1986; R\\\"udiger \\& Kichatinov 1990). More specifically, these turbulence effects may be critical for explaining the narrow structure of the observed solar torsional oscillations (Kleeorin et al.\\ 1996).  Clearly, there are several possibilities of improvement that might make the model more realistic and eventually suitable for confrontation against observations. Most important is perhaps the fact that our reference values $B_{{\\rm p}0}$ and $B_{{\\rm s}0}$ in the quenching profiles \\eqs{qp}{qs}, as well as the coefficients $q_{{\\rm p}0}$ and $q_{{\\rm s}0}$ in these profile functions, are currently kept constant. It will be more realistic to make them vary with depth, because density and turbulent rms velocity vary with depth. Another extension of the model would be to allow for the possibility that the magnetic field to be generated by a mean-field dynamo rather than relying on an imposed field.  On the more theoretical side, there is a need to further verification of the essential physics of the negative magnetic pressure effect. In particular, one must wonder why the effects of this instability have not yet seen in DNS. A likely possibility is the lack of sufficient scale separation. Only now realistic simulations are beginning to be large enough to encompass sufficiently many cells in all three directions. The importance of sufficient horizontal extent was already emphasized by Tao et al.\\ (1998), who find clear evidence of a segregation into strongly magnetized and weak magnetized regions, a phenomenon that might be closely related to that reported here. A particularly promising approach might be to generalize the direct simulations discussed in the present paper to the case with vertical density stratification such that the setup becomes similar to the mean field models that we also studied in this paper. The turbulence would then still be driven by a forcing function. Another possibility is to study this effect with turbulent convection instead of forced turbulence. This is particularly interesting, because theoretical predictions of Rogachevskii \\& Kleeorin (2007) suggest that the modification of the effective Lorentz force will be even stronger in turbulent convection."
    },
    "0910/0910.1004_arXiv.txt": {
        "abstract": "Recent observations of exoplanets by direct imaging, reveal that giant planets orbit at a few dozens to more than a hundred of AU from their central star. The question of the origin of these planets challenges the standard theories of planet formation. We propose a new way of obtaining such far planets, by outward migration of a pair of planets formed in the 10~AU region. Two giant planets in mean motion resonance in a common gap in the protoplanetary disk migrate outwards, if the inner one is significantly more massive than the outer one. Using hydrodynamical simulations, we show that their semi major axes can increase by almost one order of magnitude. In a flared disk, the pair of planets should reach an asymptotic radius. This mechanism could account for the presence of Fomalhaut~b\\,; then, a second, more massive planet, should be orbiting Fomalhaut at about 75~AU. ",
        "introduction": "Most of the known exoplanets are gaseous giants, with semi-major-axes below $2$~AU, or even as close to their parent star as $0.01$~AU. Inward planetary migration is generally considered as a necessary phenomenon to explain this distribution, because these planets should not form there. In the core-accretion model \\citep{Pollack-etal-1996}, giant planets should form beyond the ice-line (located at about $4$~AU for a solar type star). The general expectation is that planets form at distances comparable to those characterizing the orbits of the giant planets of our solar system, because further out the dynamical (and accretional) timescales are too long.  However, giant planets have been recently observed by direct imaging at a few dozens to $120$~AU from their host star \\citep{Kalas-etal-2008,Marois-etal-2008}. It has been proposed that these large semi-major-axes may be due to scattering with other giant planets after formation in the $5-20$~AU region \\citep[e.g.][]{Scharf-Menou-2009,Veras-etal-2009}. \\citet{Boley-2009}, \\citet{Clarke-2009}, and \\citet{Rafikov-2009} suggest that planet formation by gravitational instability \\citep[see][for a review]{Durisen-etal-PPV} can be effective beyond $\\sim 50$~AU, because the cooling time relative to the dynamical time becomes short, and the Toomre $Q$ parameter may be smaller in the colder, outer parts of the disk.  In this Letter, we explain how resonant interactions between two planets in a common gap can lead to outward migration, possibly explaining the presence of these planets. We outline the mechanism in Sect.~\\ref{sec:review}, and present proof-of-concept simulations in Sect.~\\ref{sec:long}. The effects of disk structure are discussed in Sect.~\\ref{sec:concept}, where we show that a pair of planets can be driven to an equilibrium radius in a flared disk. Caveats are discussed in Sect.~\\ref{sec:discussion}, and our main conclusions and applications of this mechanism, paying special attention to the case of Fomalhaut, are presented in Sect.~\\ref{sec:conclu}. ",
        "conclusions": "\\label{sec:discussion}  \\subsection{Gas accretion on the planets} \\label{sec:accretion}  A shortcoming of the mechanism presented here is that it disregards the accretion of gas onto the planets. As the pair of planets proceeds outwards, the gas proceeds inwards through the common gap. The outer planet should presumably accrete most of the incoming material, narrowing the mass difference with the inner planet, so that the torque balance could cancel or reverse at some point. Accretion of the nebular gas onto a giant planet is a complex process. It is dependent on the global structure of the flow, on the local properties of the circum-planetary disk or envelope (poorly described in numerical simulations), and on the micro-physics of the nebula, such as the opacity of the grains. Most detailed studies of gas accretion onto giant planets can't be applied to the present case, because the conditions are different from the conditions that prevail at $10$~AU, where these studies apply. This appeals for a thorough investigation of gas accretion in the conditions of outward migration in the outskirts of the disk, which is beyond the scope of this work.   \\subsection{Self-gravity of the gas and 3D effects}  The gas self-gravity may yield, depending on the equation of state of the fluid and the mass of the disk, to a vertical compression of the fluid in the vicinity of the shock \\citep[e.g.][]{Boley-Durisen-2006}. This would increase the Lindblad torque, more efficiently on the side of the most massive planet, thus it would actually enhance the mechanism of outward migration. But 3D effects apply mostly at high order resonances, close to the planet, which are depleted in our case. In addition, three dimensional semi-analytic calculations performed by \\citet{Tanaka-etal-2002} show that the angular momentum is essentially exchanged with waves that have no vertical structure. Therefore, owing to the size of the gap that our pair of planets carves in the disk, and owing to the relatively large value of the Toomre parameter (see Table~\\ref{table}), we feel confident that non-self-gravitating, two dimensional isothermal simulations capture the main features of our mechanism and provide an acceptable order of magnitude of the outward drift rate.   \\subsection{Equation of state} \\label{sec:eos}  In the above simulations, the equation of state was locally isothermal. This is realistic at large distances from the star, where the opacity is low and the cooling time is short with respect to the dynamical time. However, at $\\sim 10$~AU from the star, the cooling time can be larger than the orbital period. This can affect the aspect ratio, because of the heating of the disk by the wakes launched by the planets.  Additional simulations have been run, with the settings of Sect.~\\ref{sec:result}, and computing the full energy equation, with viscous heating $Q_+$, and no heat diffusion in the disk plane, but vertical radiative cooling $Q_-=2\\sigma_R T^4 / \\kappa\\Sigma$ (where $\\sigma_R$ is the Stefan-Boltzmann constant). The opacity $\\kappa$ is constant and tuned for the unperturbed disk to be in thermal equilibrium \\citep[see Eq.~(6) of][]{Crida09}. This gives a cooling time of $111\\,\\Omega^{-1}$ initially. The aspect ratio increases to about $0.06$ in the planets region after 100 orbits of the inner planet. This makes the migration rate after the resonance capture to decrease to $\\sim 10^{-4}$~AU/year instead of $\\sim 4\\times 10^{-4}$~AU/year.  With a realistic opacity, given by \\citet[][Appendix]{BL1994}, $T$ and $H/r$ are increased for $r \\lesssim 10$~AU, and shrink for $r\\gtrsim 15$~AU because we don't take heating from the star into account. At $10$~AU, the cooling time is then $\\sim 400\\,\\Omega^{-1}$. Migration proceeds outwards at a speed initially similar to the locally isothermal case (which was expected as $H/r$ is not much changed around $10-15$~AU), slightly accelerating while reaching the outer regions where $H/r$ is smaller.  So, it seems that migration can proceed outwards on the long range also in non locally isothermal disks, even if a self-consistent radiative disk model still has to be tested."
    },
    "0910/0910.4883_arXiv.txt": {
        "abstract": "We report on  the second \\agile{} multiwavelength campaign of the blazar 3C 454.3 during the first half of December 2007. This  campaign involved AGILE, {\\it Spitzer}, {\\it Swift}, {\\it Suzaku}, the WEBT consortium, the REM and MITSuME telescopes, offering a broad band coverage that allowed for a simultaneous sampling of the synchrotron and inverse Compton (IC) emissions. The 2-week \\agile{} monitoring was accompanied by radio to optical monitoring by WEBT and REM and by sparse observations in mid-Infrared and soft/hard X-ray energy bands performed by means of Target of Opportunity observations by {\\it Spitzer}, {\\it Swift} and {\\it Suzaku}, respectively. The source was detected with an average flux of $\\sim 250 \\times 10^{-8}$ \\phcmsec above 100 MeV, typical of its flaring states. The simultaneous optical and $\\gamma$-ray monitoring allowed us to study the time-lag associated with the variability in the two energy bands, resulting in a possible $\\lesssim$ 1-day delay  of the gamma-ray emission with respect to the optical one. From the simultaneous optical and gamma-ray fast flare detected on December 12, we can constrain the delay between the gamma-ray and optical emissions within 12 hours. Moreover, we obtain three Spectral Energy Distributions (SEDs) with simultaneous data for 2007 December 5, 13, 15, characterized by the widest multifrequency coverage. We found that a model with an external Compton  on seed photons by a standard disk and reprocessed by the Broad Line Regions does not describe in a satisfactory way the SEDs of 2007 December 5, 13 and 15. An additional contribution, possibly from the hot corona with $T = 10^{6}$ K surrounding the jet, is required to account simultaneously for the softness of the synchrotron and the hardness of the inverse Compton emissions during those epochs. ",
        "introduction": "3C 454.3 is a flat spectrum radio quasar at redshift $z=0.859$. It is one of the  brightest extragalactic radio sources with superluminal motion hosting a radio and X-ray jet. It has been observed in almost all the electromagnetic spectrum from radio up to \\gray energies; the SED has the typical double-humped shape of the blazars, the first peak occurring at mid-far infrared frequencies and the second one at MeV--GeV energies (see Ghisellini et al. 1998).  The first peak is commonly interpreted as synchrotron radiation from high energy electrons in a relativistic jet, while the second component is due to electrons scattering on soft seed photons. In the context of a simple, homogeneous scenario, the emission at the synchrotron and IC peaks is produced by the same electrons population that can self scatter the same synchrotron photons (Synchrotron Self Compton, SSC). Alternatively, the jet-electrons producing the synchrotron flux can Compton-scatter seed photons produced outside of the jet (External Compton, EC).  3C 454.3 is a highly variable blazar source. In spring 2005,  3C 454.3 experienced a strong outburst in the optical band reaching its historical maximum with $R=12.0 $ mag  (Villata et al. 2006). The exceptional event triggered observations at higher energies from Chandra (Villata et al. 2006), {\\it Swift} (Giommi et al. 2006) and INTEGRAL (Pian et al. 2006).  The available data allowed to build the spectrum of the source up to 200 keV. In particular, INTEGRAL detected  (15-18 May 2005) the source from 3 up to 200 keV in a very bright state ($\\sim  5\\times 10^{-10}$ erg cm$^{-2}$ s$^{-1}$), being almost a factor 2-3 higher than the previously observed fluxes (see Tavecchio et al. 2002). Pian et al. (2006) compared the SED in 2000 with that obtained during the 2005 outburst. They were able  to describe both observed SEDs with minimal changes in the jet power, assuming that the dissipation region (where most of the radiation is produced) was inside the Broad Line Region (BLR) in 2000 and outside of it  in 2005. On the other hand, Sikora et al. (2008) argued that X-rays and \\gray could be produced via inverse Compton scattering of near- and mid-IR photons emitted by the hot dust: a very moderate energy density of the dust radiation is sufficient to provide the dominance of the EC luminosities over the SSC ones.  In July 2007 the source woke up again in the optical band, reaching a maximum at $R=12.6$  mag (Raiteri et al. 2008b). Such an increase in the optical activity triggered observations with \\swi{} and \\agile{}.  Although still in its Science Verification Phase,  \\agile{} repointed at 3C 454.3 and detected it in high $\\gamma$-ray activity. The average \\gray flux detected by \\agile{} was the highest $\\gamma$-ray flux ever detected from this blazar, being $(280 \\pm 40) \\times 10^{-8}$ \\phcmsec (see Fig. 3 lower panel in Vercellone et al. 2008).  Ghisellini et al. (2007) reproduced the three source states in 2000, 2005, 2007  with the model proposed in Katarzynski $\\&$ Ghisellini (2007). The model  assumes that the relative importance of synchrotron and SSC luminosity with respect to the EC one is controlled by the value of the bulk Lorentz factor $\\Gamma$, which is associated to the compactness of the source.  Villata et al. (2007) and Raiteri et al. (2008b) suggested an alternative interpretation involving changes of the viewing angle of the different emitting regions of the jet.  In both cases, a strong degeneracy of parameters exists in both SSC and  EC models especially when the \\gray data are missing. This is the case of both SEDs in 2000, 2005 in which historical EGRET data were used to constrain the models.  3C 454.3 exhibited outbursts several times between 2007 and 2008 (see Vercellone et al. 2008, Tosti et al. 2008, Raiteri et al. 2008a, Raiteri et al. 2008b, Vercellone et al. 2009)  posing stringent constraints on its gamma-ray duty cycle.  This strengthened the need for simultaneous observations in different energy bands. In the case of 3C 454.3 (and other MeV blazars) it is clear that the dominant contribution in the SED comes from IR-optical bands, where the synchrotron peak lies and from both X-rays and \\gray energy range where the inverse Compton emission lies.  In this paper we present and discuss the result of a multiwavelength campaign on 3C 454.3 during a period of intense $\\gamma$-ray activity  occurred between 2007 December 1 and 16. In Section 2 we present the multiwavelength campaign, in Sect. 3 - 7 we present the \\agile, \\suzaku{}, \\swi{}, $Spitzer$, REM, WEBT and \\MITSUME{} observations and data analysis; in Sect.\\  8 we analyse the $\\gamma$-optical correlation and present broad-band SEDs built with simultaneous data, discussing in details how they are modelled in the framework of SSC and EC scenarios. Throughout this paper the photon indexes are parametrized as $N(E)\\propto E^{-\\Gamma}$ (\\phcmsec keV$^{-1}$ or MeV$^{-1}$). The uncertainies are given at 1-$\\sigma$ level, unless otherwise stated. ",
        "conclusions": "We reported in this paper the main results of a multifrequency campaign on the blazar 3C 454.3 performed in December 2007. The source was simultaneously observed in mid-Infrared, optical, X-ray and \\gray energy bands, which provided us with a wide dataset aimed to study the correlation between the emission properties at lower and higher frequencies. We summarize below the major results.  \\begin{itemize} \\item The \\gray emission from 3C 454.3 shows variations on a daily time scale. \\item The simultaneous monitoring of the source in the optical and \\gray energy bands allowed us to determine a possible $\\lesssim$ 1 day delay of the \\gray emission with respect to the optical one. \\item The extraordinary optical activity (lasting less than 3 hours), observed on December 12 has a counterpart in the \\gray data. A possible delay between the \\gray emission and the optical one is constrained within 12 hours. \\item We found that a leptonic model with an External Compton on seed photons from disk and BLR does not succeed to account for both the ``hardeness'' of the \\gray spectrum and the ``softness'' of the Synchrotron emission, requiring  an additional component. We argued that a possible candidate for it is  the hot Corona ($T \\sim 10^{6}$ K) surrounding the disk. \\end{itemize}"
    },
    "0910/0910.1967_arXiv.txt": {
        "abstract": "A large sample of low surface brightness (LSB) disk galaxies is selected from SDSS with B-band central surface brightness $\\mu_0$(B) from 22 to 24.5 mag arcsec$^{-2}$. Some of their properties are studied, such as magnitudes, surface brightness, scalelengths, colors, metallicities, stellar populations, stellar masses and multiwavelength SEDs from UV to IR etc. These properties of LSB galaxies have been compared with those of the galaxies with higher surface brightnesses. Then we check the variations of these properties following surface brightness. ",
        "introduction": "Low Surface Brightness Galaxies (LSBGs) are important populations in galaxy field. However, their  contributions to galaxy population have been underestimated for a long time since they are hard to find owing to their faintness compared  with the night sky.  An initial quantitative study was made by Freeman (1970), who noticed that the central surface brightness of their 28 out of 36 disc  galaxies fell within a rather  narrow range, $\\mu_0$(B)=21.65$\\pm$0.3 mag arcsec$^{-2}$. This could be caused by selection effects (Disney 1976).  Since then, many efforts have been made to search for large number of LSBGs  from surveys (Bothun \\& Impey 1997; Impey \\& Bothun 1997;  Zhong et al. 2008 and references therein).  One of the important ones is the APM survey, i.e., Impey et al. adopted the Automated Plate Measuring (APM) mechanism to scan UK Schmidt plates and  discovered 693 LSBGs  which forms the most extensive catalog of LSBGs to that date (Impey et al. 1996). O'Neil et al. (1997, in ``Texas survey\")  firstly found the red LSBG populations. Monnier Ragaigne et al. (2003) selected a sample of about 3800 LSBGs from the all-sky near-infrared 2MASS survey, and then made HI observations for a sub-sample.  The modern digital sky survey, such as the Sloan Digital Sky Survey (SDSS),  certainly provide a wonderful chance for us to find much more LSBGs.  Kniazev et al. (2004) developed a method to search for LSBGs from SDSS images and used the APM sample to test their method. They recover 87 same objects and 42 new LSBGs. Recently, our group (Zhong et al. 2008) successfully select a large sample of LSBGs from SDSS-DR4 main galaxy sample, which includes 12,282 nearly face-on disk galaxies with low surface brightness $\\mu_0$(B)$\\geq$ 22 mag arcsec$^{-2}$, and another 18,051 high surface brightness galaxies (HSBGs) with $\\mu_0$(B)$<$22 mag arcsec$^{-2}$ are also selected for comparisons. This will be a very efficient sample  to study the properties of LSBGs, and to check the variation of these properties with surface brightnesses if there exists.  We have studied some properties of this large sample. Zhong et al. (2008) presented the sample selection criteria, and their magnitudes, surface brightness, scalelengths and colors etc. The metallicities (Liang et al. 2009, in preparation) and  stellar populations (Chen et al. 2009, in preparation) are also studied from optical spectra, i.e. the emission lines, continua and absorption lines. In addition, the modern digital sky surveys have been made in  wide wavelength ranges, such as GALEX-UV, SDSS-optical, 2MASS-NIR, IRAS-IR etc. These allow us to construct the multi-wavelength SEDs for our sample galaxies  (Gao et al. 2009, in preparation), which are very useful and efficient to study their star formation history. In this paper, we summarize the interesting results we obtained for this large sample of  galaxies, which will be presented in detail in the  series work mentioned above. Especially, we try to  check the varying trends of their properties following surface brightness. A cosmological model with $H_0$ = 70 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_M$ = 0.3  and $\\Omega_\\lambda$ = 0.7 is adopted in this paper. ",
        "conclusions": ""
    },
    "0910/0910.1010_arXiv.txt": {
        "abstract": "We present timing and spectral analysis of RXTE-PCA observations of SMC X-1 between January 1996 and December 2003. From observations around 30 August 1996 with a time span of $\\sim 6$ days, we obtain a precise timing solution for the source and resolve the eccentricity as 0.00089(6). We find an orbital decay rate of $\\dot P_{orb}/P_{orb} =-3.402(7) \\times 10^{-6}$ yr$^{-1}$ which is close to the previous results. Using our timing analysis and the previous studies, we construct a $\\sim 30$ year long pulse period history of the source. We show that frequency derivative shows long (i.e. more than a few years) and short (i.e. order of days) term fluctuations. From the spectral analysis, we found that all spectral parameters except Hydrogen column density showed no significant variation with time and X-ray flux. Hydrogen column density is found to be higher as X-ray flux gets lower. This may be due to the increase in soft absorption when the pulsar is partially obscured as in Her X-1 or may just be an artifact of the tail of a soft excess in energy spectrum.  {\\bf{Keywords:}} X-rays: binaries, pulsars,  individual: SMC X-1 , stars: neutron, accretion, accretion discs ",
        "introduction": "The high mass X-ray binary (HMXB) system SMC X-1/Sk 160 consists of a neutron star with a mass of $\\sim1.06M_{\\odot}$ (van der Meer et al. 2007) and a spin period of 0.71s (Lucke et al.1976), accreting from the B0I star Sk 160 with a mass of  $\\sim17.2M_{\\odot}$ (Reynolds et al. 1993). Orbital period of the system is $\\sim3.9$ days (Schreier et al. 1972). The system has been observed with several observatories since it was discovered during a rocket flight (Price et al. 1971). It is the only discovered HMXB with a supergiant companion in SMC so far (Galache et al. 2008).  From the pulse timing studies, the source was observed to be spinning up since it was discovered (Wojdowski et al. 1998). Levine et al. (1993) and Wojdowski et al (1998) found an orbital period decay in the system with $P_{orb}/\\dot{P}_{orb}$ of $\\sim 3.4\\times 10^{-6}$yr$^{-1}$.  SMC X-1 also exhibits super-orbital X-ray flux variations like the 35 day cycle of Her X-1 (Gruber $\\&$ Rothschild, 1984). Average super-orbital period of the source is $\\sim$ 55 days (Wojdowski et al. 1998; Trowbridge et al. 2007). The super-orbital X-ray flux variations of SMC X-1 are thought to be due to the precession of a warped accretion disk (Wojdowski et al. 1998).  In this paper, we present timing and spectral analysis of RXTE (Rossi X-ray Timing Explorer) observations of SMC X-1 between January 1996 and December 2003. In the next section, we give brief information about instruments and observations. In Section 3, we present the results of timing analysis, including pulse period history and timing solution of the source. In Section 4, X-ray spectral analysis of the source is presented. ",
        "conclusions": "In this paper, we presented timing and spectral analysis of RXTE-PCA observations of SMC X-1 with a time span of about 8 years. From timing analysis, we revised timing solution of the source and resolved an eccentricity value. We also confirmed the orbital decay reported before. Our timing analysis helped us to construct a $\\sim 30$ year long pulse period history of the source (see Figure 3).  From the spectral analysis, we found that all spectral parameters except Hydrogen column density showed no significant variation with time and X-ray flux.  Figure 1 demonstrates that the eccentric orbit model is a better fit compared to the circular orbital model. The timing solution of the source revealed that the binary orbit has an eccentricity of 0.00089(6). Wojdowski et al. (1998) had found a circular orbit solution and Levine et al. (1993) had only found an upper limit of 0.00004 for eccentricity. Our present value is more than 20 times greater than the upper limit found by Levine et al. 1993. This significant eccentricity value should be verified using future monitoring observations.  From timing analysis, we obtained a new orbital epoch of SMC X-1 from RXTE observations. Using this new epoch and previous results (Wojdowski et al. 1998), we found that there is an orbital decay with $\\dot{P_{orbit}}/P_{orbit}$ of $-3.402(7)\\times10^{-6}$ yr$^{-1}$. This value is similar to the values found by Wojdowski et al. (1998) and Levine et al. (1993). Levine et al. (1993) proposed that the major cause of change in the orbital period is tidal interactions as for the case in Cen X-3 (Kelley et al. 1983), and LMC X-4 (Levine et al. 2000). SMC X-1 is unlike Her X-1 for which the mass accretion were thought to be primarily related to the orbital period decay (Deeter et al. 1991).  From Figure 3, it is seen that SMC X-1 spins up continuosly for $\\sim 30$ years. We found that average long term spin-up rate of the source between MJD 50093.048 and MJD 52987.640 (whole time span of our analysis) is $2.65343816(7)\\times 10^{-11}$ Hz s$^{-1}$.  Spin-up rate of a source with a persistent accretion disk can be expressed as  \\begin{equation} \\dot{\\nu}\\simeq 2.2\\times10^{-12}\\mu_{30}^{2/7}m_x^{-3/7}R_6^{6/7}I_{45}^{-1}L_{37}^{6/7} {\\rm{Hz\\,s}}^{-1}, \\end{equation}  \\noindent{where $\\dot{\\nu}$ is first time derivative of the spin frequency, $\\mu_{30}$ is the magnetic moment of the neutron star in units of $10^{30}$ Gauss cm$^{3}$, $m_x$ is the mass of the neutron star in units of solar mass, $R_6$ is the radius of the neutron star in units of cm, $I_{45}$ is the moment of inertia of the neutron star in terms of $10^{45}$ g cm$^2$, and $L_{37}$ is the luminosity of the neutron star in  units of $10^{37}$ erg s$^{-1}$ (Ghosh\\& Lamb, 1979). Using a distance value of 61kpc (Keller\\& Wood 2006) and a flux of $1.2\\times 10^{-9}$ ergs cm$^{-2}$ s$^{-1}$ (see Figure 4), luminosity of SMC X-1 is $\\sim 5.5\\times 10^{38}$ erg s$^{-1}$. Assuming magnetic moment of $10^{30}$ Gauss cm$^3$, mass of $1.4M_{\\odot}$, radius of $10^{6}$ cm, moment of inertia of $10^{45}$ g cm$^{2}$, we find $\\dot{\\nu}$ to be $\\sim 5.9\\times 10^{-11}$ Hz s$^{-1}$. This value is of the order of the observed spin-up rate value of the source.}  From Figure 3 and Table 4, it is evident that the long term spin-up rate of SMC X-1 varies within a factor of 2. From Equation 3, this variation may -in principle- be related to a change in X-ray luminosity of the source. However, previous X-ray observations of the source were performed by different X-ray observatories and/or in different energy bands. These observations are sparsely distributed and average X-ray flux of these observations are not likely to represent an average X-ray flux for the intervals for which the pulse frequency derivatives were found. Moreover, X-ray flux variations are also affected by the super-orbital cycle, period of which is not as stable as that of Her X-1 (Clarkson et al. 2003). Therefore, it is not possible to test whether variation in long term spin-up rate is related to a change in X-ray luminosity or not.  From Table 1, spin-up rate obtained from $\\sim 6$days long observation around MJD 50324.7 is $3.279(5)\\times 10^{-11}$ Hz s$^{-1}$. This is about 20$\\%$ greater than the long term spin-up rate between MJD 50093 and 52988. Assuming that the pulse frequency variations can be explained in terms of random walk model in pulse frequency (Baykal \\& Ogelman 1993), one can express the pulse frequency derivative variations as $<\\Delta\\dot{\\Omega}^{1/2}>=\\sqrt{{S}\\over{t}}$ where S is the noise strength and t is the time interval at which the pulse frequency derivatives are observed. Therefore it is natural to observe high pulse frequency derivative fluctuations at shorter time scales. It is important to note that we refined pulse period values fitting arrival times each obtained from $\\sim 200$s long intervals, therefore we had to have at least a few arrival times to obtain spin period accurately. So, our shortest timescale to obtain pulse periods is of the order of a single RXTE observation which is $\\sim 1-2$ksec long.  From the spectral analysis, we found that all of the spectral parameters except Hydrogen column density did not vary significantly. Hydrogen column density was found to be higher as X-ray flux gets lower.  Increase in Hydrogen column density with a decrease in X-ray flux is also observed in Her X-1 (Inam\\& Baykal, 2005) and may be due the fact that soft absorption becomes stronger whenever there is a partial obscuration of the neutron star due to the X-ray eclipses and warping of the accretion disk. The increase in Hydrogen column density may also be an artifact of the simple absorbed power law model. Paul et al. (2002) showed that SMC X-1 has soft excess especially for energies $\\aplt 3$ keV. Although we used photon energies greater than 3keV in our analysis, tail of a soft spectral compenent may affect low energies in our analysis and as a result we might have misidentified corresponding changes in energy spectrum as  variations in Hydrogen column density parameter."
    },
    "0910/0910.3015_arXiv.txt": {
        "abstract": "We explore for the first time the method of cross-correlation of radio synchrotron emission and tracers of large-scale structure in order to detect the diffuse IGM/WHIM. We performed a cross-correlation of a 34$^{\\circ}\\times34^{\\circ}$ area of 2MASS galaxies for two redshift slices (0.03 $<$ z $<$ 0.04 and 0.06 $<$ z $<$ 0.07) with the corresponding region of the 1.4 GHz Bonn survey. For this analysis, we assumed that the synchrotron surface brightness is linearly proportional to surface density of galaxies.  We also sampled the cross-correlation function using 24 distant fields of the same size from the Bonn survey, to better assess the noise properties. Though we obtained a null result, we found that by adding a signal weighted by the 2MASS image with a filament (peak) surface brightness of 1 (7)~mK and 7 (49)~mK would produce a 3$\\sigma$ positive correlation for the 0.03 $<$ z $<$ 0.04 and 0.06 $<$ z $<$ 0.07 redshift slices respectively. . These detection thresholds correspond to minimum energy magnetic fields as low as 0.2~$\\mu$G, close to some theoretical expectations for filament field values.  This injected signal is also below the rms noise of the Bonn survey, and demonstrates the power of this technique and its utility for upcoming sensitive continuum surveys such as GALFACTS at Arecibo and those planned with the Murchison Widefield Array (MWA). ",
        "introduction": "\\vspace{-0.1in} In the current cosmological paradigm, $\\approx$95\\% of the mass/energy density of the universe is composed of dark energy and dark matter, both of which have yet to be directly detected. The remaining 5\\% are ordinary baryons, which at redshifts of z $>$ 2 are fully accounted for based on Ly$\\alpha$ forest observations of the photoionised intergalactic medium (IGM) and ordinary galaxies (e.g. Rauche et al. 1997; Weinberg et al. 1997; Schaye 2001). In the current epoch though, roughly half of these baryons are missing. Simulations suggest that the collapsing diffuse IGM was shock-heated and now resides in filaments as T$\\sim$10$^{5}$ - 10$^{7}$~K WHIM, where it is practically invisible at most wavelengths (Cen \\& Ostriker 1999, 2006; Dav{\\'e} et al. 2001). A few tentative absorption detections of the coolest WHIM components using OVII and OVIII have been reported, but they give no information on the spatial distribution. However, shocks from infall into and along the filamentary structures between clusters are now widely expected to generate relativistic plasmas which track the distribution of the WHIM (Keshet et al. 2004; Pfrommer et al. 2006; Ryu et al. 2008; Skillman et al. 2008); indeed, merger/accretion shocks are seen as polarized radio sources (the ``peripheral relics\") at the edges of the dense X-ray gas in clusters.  This radio emission also has the potential for probing into the lower density regions further from cluster cores. When such features are detected, they can be used to set limits on the pressure of the (invisible) thermal gas, delineate shock structures, and illuminate large scale magnetic fields. The next generation of low frequency radio telescopes (e.g., LOFAR, MWA, LWA, and eventually SKA), as well as current deep continuum surveys (e.g., GALFACTS\\footnote{http://www.ucalgary.ca/ras/GALFACTS/} at Arecibo), have the potential to map this steep-spectrum, ``cosmic-web\" of synchrotron emission (Battaglia et al. 2009).  However, there are two major problems that will challenge future radio surveys in their efforts to detect this emission. 1) Disentangling the extragalactic signal from diffuse Galactic emission that has higher surface brightness and similar angular scale; 2) Determining the redshift of the diffuse extragalactic signal (e.g., Brown \\& Rudnick 2009).  One possible solution is to cross-correlate large field-of-view synchrotron maps with optical/IR tracers of large-scale structure (LSS; Keshet et al. 2004). The cross-correlation function (CCF) has long been used as means of detecting faint emission, often below the image sensitivity, in CMB maps such as the SZ and late ISW effects (e.g., Boughn \\& Crittenden 2004). Cross-correlation can detect signals that are buried deep in the noise when the noise is not spatially correlated with the tracer image. In this case, the Galactic synchrotron ``noise\" is not expected to be spatially correlated with the large-scale distribution of galaxies, while the synchrotron emitting relativistic plasma associated with the WHIM is correlated (Battaglia et al. 2009). Attempts to detect the WHIM with cross-correlation using the SZ effect has been tried with WMAP and 2MASS data (e.g., Hernandez-Monteagudo et al. 2004; Hansen et al. 2005; Cao et al. 2006), without success. In this paper we outline the first attempt to detect a correlation between synchrotron radio emission and large-scale structure, and demonstrate the power of cross-correlation by detecting simulated emission added into the synchrotron images. In this paper, we assume $H_{o}=70$, $\\Omega_{\\Lambda}=0.7$, $\\Omega_{M}=0.3$.  \\vspace{-0.3in} ",
        "conclusions": "We have explored for the first time the method of cross-correlation of diffuse radio emission and tracers of LSS in order to detect the synchrotron radiation associated with the diffuse baryons in the WHIM. Filtering out the large-scale Galactic emission is critical for reducing the CC noise. We demonstrated that, if the filamentary radio signal is traced perfectly by the distribution of galaxies, emission significantly below the sensitivity of the Bonn survey can be statistically detected. Our detection thresholds are already at a level expected in some models, and dramatic improvements using new surveys will be possible in the future.  Further refinements to the CC technique, including subtraction of point sources which also trace the LSS, will also improve its use as a detection tool for the diffuse baryons in filaments. The assumption of perfect correspondence between the distribution of galaxies and IGM radio emission is clearly an oversimplification, and further experimentation with numerical simulations is needed.  Partial support for this work at the University of Minnesota comes from the U.S. National Science Foundation grant AST~0607674. We acknowledge the use of NASA's SkyView facility located at NASA Goddard Space Flight Center. SB also acknowledges support from a Doctoral Dissertation Fellowship from the Graduate School of the University of Minnesota."
    },
    "0910/0910.4403_arXiv.txt": {
        "abstract": "We combine near-infrared (2MASS) and mid-infrared ({\\it Spitzer}-IRAC) photometry to characterize the IR extinction law (1.2-8 $\\mu$m) over nearly 150$^\\circ$ of contiguous Milky Way midplane longitude. The relative extinctions in 5 passbands across these wavelength and longitude ranges are derived by calculating color excess ratios for G and K giant red clump stars in contiguous midplane regions and deriving the wavelength dependence of extinction in each one. Strong, monotonic variations in the extinction law shape are found as a function of angle from the Galactic center, symmetric on either side of it. These longitudinal variations persist even when dense interstellar regions, known a priori to have a shallower extinction curve, are removed. The increasingly steep extinction curves towards the outer Galaxy indicate a steady decrease in the absolute-to-selective extinction ratio ($R_V$) and in the mean dust grain size at greater Galactocentric angles. We note an increasing strength of the 8\\mic extinction inflection at high Galactocentric angles and, using theoretical dust models, show that this behavior is consistent with the trend in $R_V$. Along several lines of sight where the solution is most feasible, $A_\\lambda$/$A_{Ks}$ as a function of Galactic radius ($R_{GC}$) is estimated and shown to have a Galactic radial dependence. Our analyses suggest that the observed relationship between extinction curve shape and Galactic longitude is due to an intrinsic dependence of the extinction law on Galactocentric radius. ",
        "introduction": "Interstellar extinction presents one of the longest-standing challenges to observational astronomy. Extensive work has been done to understand the complex and highly-variable ultraviolet and optical portions of the extinction curve (see, e.g., reviews by \\citealt{Mathis_90_dustreview} and \\citealt{Whittet_03_dust}); only recently, however, has comprehensive characterization of the infrared (IR) extinction law become possible, through the influx of near- and mid-infrared (NIR, MIR) data provided by large-area surveys like 2MASS, UKIDSS, and GLIMPSE I/II/3D. Not only does this characterization enable more accurate study of sources extinguished by dust, but reliable derivation of the wavelength-dependent IR extinction behavior also facilitates identification and modeling of the physical properties (e.g., chemical composition, grain size, crystallization fraction) of the interstellar dust itself.  Many dust absorption studies to date have converged to a treatment of the infrared extinction law in the diffuse interstellar medium (ISM) as nearly constant and universal. The NIR extinction curve ($\\lambda$ $\\lesssim$ 1--4 $\\mu$m) is often modeled as a power law ($A_\\lambda\\propto\\lambda^{-\\beta}$), with $\\beta$ ranging between 1.6 and 1.8 \\citep[e.g.,][and references therein]{RiekeLeb_85_extlaw,Draine_03_dust}, but values up to 2 have been reported \\citep{Nishi_06_extlaw,Nishi_09_extlaw}. At longer wavelengths ($\\lambda$ $>$ 4 $\\mu$m), however, $A_\\lambda$ deviates from this relationship and appears to ``flatten out'' (become grayer) before increasing towards the peak of the 9.7\\mic silicate feature. The absolute extinction ratios $A_{\\lambda_0}$/$A_{Ks}$ used to normalize the derived extinction curves are generally determined for the diffuse ISM using reddened background standard candles for which intrinsic colors can be assumed---e.g., red clump stars \\citep{Indeb_05_extlaw,Nishi_09_extlaw} or RGB-tip or low-mass-loss AGB stars \\citep{Jiang_03_extlaw,Jiang_06_7-15ext}.  More recent work has demonstrated that in regions of dense ISM, such as dark cloud cores or star-forming regions, the MIR $A_\\lambda$ curve becomes even shallower than in diffuse material, likely due to dust grain growth through grain coagulation or ice mantle coating \\citep[e.g.,][]{WeinDraine_01_dustsize,Moore_05_extlaw,RZ_07_extlaw,Flaherty_07_extlaw,McClure_08_extlaw,Chapman_08_extlaw}. The absolute ratios $A_{\\lambda_0}$/$A_{Ks}$ have been determined in these regions using either stars assumed to be in the background \\citep{Flaherty_07_extlaw} or associated gas emission---e.g., hydrogen recombination spectra \\citep{Lutz_96_galcenter,Moore_05_extlaw} or H$_2$ rotational and rovibrational lines \\citep{Rosenthal_00_H2extlaw}.  Infrared extinction behavior as a function of density, beyond a simple ``dense'' vs. ``diffuse'' dichotomy, has only recently begun to be studied in specific cloud cores \\citep{Chapman_08_extlaw,McClure_08_extlaw}. We address the effects of ISM density on the extinction law here as well, but the present work is also the first to examine extinction behavior with respect to overall Galactic ISM density or grain-size gradients.  Given the differences in the extinction law behavior in dense and diffuse interstellar environments, one may expect the Galactic extinction law to vary substantially through the widely changing environments across the disk. Yet the vast majority of infrared extinction law studies, including those listed above, use no more than a few specific regions on the sky ($\\lesssim$10 deg$^2$) to define the extinction law and/or are limited to measuring relative $A_\\lambda$ variations in environmental extremes, such as dark clouds. Moreover, most efforts are concentrated in the inner Galaxy, where both starcounts and the overall extinction are higher, but where the potential effects of Galactic-scale gradients (such as ISM metallicity or density) cannot be easily identified. Thus, there is little reason for the ``universal'' extinction law derived from these observations to be in fact applicable to the Galaxy as a whole.  In contrast to previous studies, this work tests both the very assumption of a universal diffuse extinction law and the pertinence of a simple bimodal {\\it dense} versus {\\it diffuse} description of the ISM with regard to extinction. Our collected databases allow us to determine coherent patterns of extinction behavior by measuring the IR extinction law in a consistent way along many different sightlines spread over $\\sim$150$^{\\circ}$ of the Galactic disk. \\S\\ref{sec:data} of this paper contains a description of the data we use in this new, comprehensive exploration of extinction behavior. \\S\\ref{sec:measure_ext} explains our methodology for determining $A_\\lambda$/$A_{Ks}$, and \\S\\ref{sec:density} examines the question of density as the sole driver of extinction variations. In \\S\\ref{sec:discuss}, we compare our results to previous observational and theoretical studies and discuss their similarities and differences with our findings. We also address the implications of the observed variations in the 8\\mic extinction upturn and the feasibility of measuring extinction behavior as a function of Galactocentric radius.  Finally, \\S\\ref{sec:conc} summarizes our findings and conclusions. ",
        "conclusions": "\\label{sec:conc}  Using data from 2MASS \\citep{Skrutskie_06_2mass} and three extensive {\\it Spitzer}-IRAC surveys \\citep[][PID 20499, and PID 40791]{Benjamin_03_glimpse}, we have studied the interstellar relative extinction as a function of wavelength through the near- and mid-infrared (1.2-8 $\\mu$m) for contiguous lines of sight covering $\\sim$150$^\\circ$ of the Galactic disk. Using G and K spectral type red clump (RC) stars, we have calculated color excess ratios and derived the wavelength-dependent extinction behavior as a function of angle from the Galactic center. We find the IR extinction law becomes increasingly steep as the Galactocentric angle increases, with identical behavior between $l$ $<$ 180$^\\circ$ and $l$ $>$ 180$^\\circ$.  Aware that dense molecular clouds scattered throughout the disk have recently been shown to exhibit MIR extinction behavior quite different from that of the diffuse ISM \\citep[e.g.,][]{Flaherty_07_extlaw,RZ_07_extlaw,McClure_08_extlaw,Chapman_08_extlaw}, we culled from our sample those stars spatially associated with lines of sight having detectable $^{13}$CO ($J$=1$\\rightarrow$0) emission \\citep[][]{Jackson_06_GRS}, here used as an indicator of high ISM density ($n$(H) $\\gtrsim$ 1--5 cm$^{-3}$).  After doing so, we find that the derived $A_\\lambda$/$A_{Ks}$ trend with Galactic angle persists, in what should be diffuse ISM only. This behavior suggests that, apart from high density molecular clouds, at least one secondary factor that varies throughout the Galactic disk is at work, such as grain size, composition, or crystalline fraction. As these factors themselves depend on local ISM conditions (including density), and that even within the diffuse ISM there are likely to be density gradients, we see clearly that the previous simple, bimodal division of the ISM into ``dense'' and ``diffuse'' environments is not adequate to characterize the displayed variation in extinction behavior.  We find close matches to our empirical extinction curves (as derived from our diffuse-only and total RC samples) in the diffuse-ISM study of \\citet{Indeb_05_extlaw} and the lowest-density cases of \\citet{Chapman_08_extlaw}, but not in the work of \\citet{Lutz_96_galcenter}, \\citet{Jiang_06_7-15ext}, and \\citet{Flaherty_07_extlaw}, who focus on regions of known high dust density. Our results are consistent with several of the theoretical dust models presented by \\citet{WeinDraine_01_dustsize}, and we use this comparison to characterize the dust at different positions in the Galaxy. We find that the best-matched $R_V$ decreases (from 5.5 to 3.1) at increasingly greater angles from the Galactic Center, and the input size distributions of the best-fit models suggest a decrease in mean grain size (from super- to sub-micron) towards the outer Galaxy.  We characterize the extinction ``upturn'' at $\\sim$8 $\\mu$m in each field's extinction law with the ratio $A_{[5.8\\mu]}$/$A_{[8\\mu]}$, and we find a clear increase in the strength of this inflection at greater Galactocentric angles. The 9.7\\mic silicate absorption, to which we attribute this increased relative extinction, is strengthened by certain changes in dust grain characteristics (e.g., to a smaller mean size and/or lesser crystalline fraction); because these changes are consistent with the trends between Galactic longitude and dust grain property derived from the model comparisons above, we interpret the 8\\mic extinction inflection behavior as further evidence for the varying nature of ISM dust grains from the inner to the outer Galaxy  We consider the problem of calculating extinction curves as a function of true Galactic radius $R_{GC}$. This is the ideal, more natural approach to parameterizing the factors that cause the observed variations in extinction behavior, but assigning the dust spread out along the line of sight to a particular $R_{GC}$ is difficult to do. To make a first attempt, we select three lines of sight for which we can make good approximations of the dust's typical $R_{GC}$ (at $\\sim$5.5, $\\sim$6.8, and $\\sim$10.4 kpc), and we find that the trend in extinction behavior strongly supports a decrease in mean dust grain size (and in $R_V$) at greater Galactic radii.  Careful measurements of the extinction in the outer part of the Galaxy are difficult to make because of the smaller numbers of observed stars and the weaker extinction to those stars. Nevertheless, large-scale surveys of the outer disk, such as the approved GLIMPSE-360 Cycle 6 {\\it Spitzer}-IRAC survey, will provide stellar photometry (at \\ione and \\itwo) of the most distant, reddened stars to the edge of the disk. A better understanding of the Milky Way extinction law as a function of Galactic position is crucial not only to a complete description of the varying dust grain characteristics but also to photometric and spectroscopic corrections made for observations of sources within or beyond our Galaxy that are affected by dust reddening and extinction."
    },
    "0910/0910.5775_arXiv.txt": {
        "abstract": "Ten years ago our team completed the Hubble Space Telescope Key Project on the extragalactic distance scale. Cepheids were detected in some 25 galaxies and used to calibrate four secondary distance indicators that reach out into the expansion field beyond the noise of galaxy peculiar velocities. The result was H$_0$ = 72 $\\pm$ 8 km s$^{-1}$ Mpc$^{-1}$ and put an end to galaxy distances uncertain by a factor of two. This work has been awarded the Gruber Prize in Cosmology for 2009. \\endabstract ",
        "introduction": "\\noindent Our story begins in the mid-1980s with community-wide discussions on Key Projects for the NASA/ESA Hubble Space Telescope. The extragalactic distance scale was the perfect example of a project that would be awarded a generous allocation of time in order that vital scientific goals would be accomplished even if the lifetime of the telescope proved to be short. As Marc Aaronson (Figure 1), the original principal investigator of the project, said in his Pierce Prize Lecture in 1985, ``The distance scale path has been a long and torturous one, but with the imminent launch of HST there is good reason to believe that the end is finally in sight.\"  \\begin{figure}[h]  \\includegraphics[width=7.5cm]{fig1.jpg} \\caption{Jeremy Mould (left) and Marc Aaronson (right) at the van Biesbroeck Prize award ceremony in 1981.} \\end{figure}  Marc Aaronson died tragically in an accident in 1987, having written a successful proposal for the Key Project, a project designed to shrink the scatter shown in Figure 2 to 10\\% and put an end to 60 years of debate, commencing with Hubble's estimates in 1929.  \\begin{figure}[ht] \\includegraphics[width=15cm]{fig2.jpg} \\caption{The scatter in estimates of the Hubble Constant showed no sign of convergence in the 1970s and 1980s. On the right is a schematic distance scale by G. de Vaucouleurs, one of the protagonists in the distance scale controversy.} \\end{figure}  The principal reason for the uncertainty in H$_0$ is evident in Figure 3. Large scale structure is seen out to distances of 100 Mpc. Ground-based Cepheid distances, however, extended to only 4 Mpc with a ``twilight zone\" beyond (Figure 4).  \\begin{figure}[ht] \\includegraphics[width=10cm]{fig3.jpg} \\caption{Colour coded redshifts show the large scale structure in the nearby Universe (galactic coordinates).} \\end{figure}  The Key Project's solution to the twilight zone problem was to map Cepheids out to 20 Mpc and calibrate secondary distance indicators within this volume. The secondary indicators would extend the distance scale out to 100 Mpc.  \\begin{figure}[ht] \\includegraphics[width=15cm]{fig4.jpg} \\caption{A schematic distance scale by other protagonists of the controversy, A. Sandage \\& G. Tammann.} \\end{figure}  An important tenet of the Key Project was to exploit the redundancy of distance indicators, as shown in Figure 5, especially the four secondary distance indicators, the Tully-Fisher relation, surface brightness fluctuations, supernovae, and the fundamental plane.  \\begin{figure}[ht] \\includegraphics[width=10cm]{fig5.jpg} \\caption{The schematic distance scale in the Key Project Committee report.} \\end{figure} ",
        "conclusions": ""
    },
    "0910/0910.2155_arXiv.txt": {
        "abstract": "New results on the B[e] star HD87643 are presented here. They were obtained with a wide range of different instruments, from wide-field imaging with the WFI camera, high resolution spectroscopy with the FEROS instrument, high angular resolution imaging with the adaptive optics camera NACO, to the highest angular resolution available with AMBER on the VLTI. We report the detection of a companion to HD87643 with AMBER, subsequently confirmed in the NACO data. Implications of that discovery to some of the previously difficult-to-understand data-sets are then presented. ",
        "introduction": "B[e] stars (or ``stars with the B[e] phenomenon'') have little similarities with the classical Be stars. They do show permitted emission lines from Hydrogen and metallic elements, and they have an infrared excess. However, they also exhibit forbidden emission lines, and their infrared excess is due to hot dust ($T\\approx 1300-1500$K) instead of heated plasma. Another difference is that the evolutionary status of B[e] stars is unclear. Some of them show characteristics typical of evolved objects (e.g. supergiant B[e] stars or SgB[e]), while others show characteristics of young stellar objects (Herbig B[e] stars are an example). Hence, ``stars with the B[e] phenomenon'' do not form a homogeneous class of objects. They have sometimes a poorly known distance, making their evolutionary status highly controversial. As a consequence, many B[e] stars are ``classified'' as ``unclassified B[e] stars'' (or UnclB[e]). One example is the star HD87643, located in the direction of the Carina arm, which has been classified as SgB[e] due to its putative distance $\\geq1$\\,kpc, and which, at the same time, shows variability typical of a young stellar object. It shows one of the most extreme infrared excess of all B[e] stars and, hence, needs to be further investigated. We have observed this star with a variety of techniques to try to fix its properties and evolutionary status. These techniques range from high-resolution spectroscopy to high angular resolution imaging. We present here the high angular resolution discovery images of a companion to HD87643, as well as a larger scale image of arcs in its surrounding nebula, that might be related to the binary star. These results were presented in details in \\citet{Millour09}. We will also present a tentative new interpretation of previously published spectro-astrometric measurements \\citep[from][that failed to detect the binary]{2006MNRAS.367..737B}, having in mind, now, that HD87643 is indeed a binary system. ",
        "conclusions": "We presented high angular resolution AMBER and NACO images as well as a high dynamic range WFI image of HD87643. AMBER and NACO images undoubtedly reveal a binary companion to the main star, while the WFI one suggest a high-eccentricity orbit for the binary. Separating the spectra of the different components in the system using the AMBER and MIDI data, it was possible to infer that the main source is likely a hot star encircled by a 6\\,AU hot dust envelope, that the secondary is embedded in a compact cocoon of dust, and that a circumbinary envelope holds most of the dust silicate emission in the system. Finally, we propose a new interpretation of literature spectro-astrometric data, in which both stellar components emit H$\\alpha$.  The global view of the system has been completely changed by our discovery of a companion to HD87643. The system might resemble in fact the following: a hot star encircled by a dusty disk, whose inner rim is seen in the AMBER image (a Herbig star?), a T Tauri companion star, and a dusty circumbinary envelope. All this suggests a much younger evolutionary status of HD87643 and, hence, a much closer distance than previously thought. A monitoring campaing of the binary throughtout its orbit would enable us to set accurately and definitively the distance of the system and, hence, its evolutionary status."
    },
    "0910/0910.0150_arXiv.txt": {
        "abstract": "{} { Emission from the 6.7 GHz methanol maser transition is very strong, is relatively stable, has small internal motions, and is observed toward numerous massive star-forming regions in the Galaxy. Our goal is to perform high-precision astrometry using this maser transition to obtain accurate distances to their host regions.} {Eight strong masers were observed during five epochs of VLBI observations with the European VLBI Network between 2006 June, and 2008 March.} {We report trigonometric parallaxes for five star-forming regions, with accuracies as good as $\\sim22~\\mathrm{\\mu}$as.  Distances to these sources are $2.57^{+0.34}_{-0.27}$~kpc for ON\\,1, $0.776^{+0.104}_{-0.083}$ kpc for L\\,1206, $0.929^{+0.034}_{-0.033}$ kpc for L\\,1287, $2.38^{+0.13}_{-0.12}$ kpc for NGC\\,281-W, and $1.59^{+0.07}_{-0.06}$ kpc for S\\,255. The distances and proper motions yield the full space motions of the star-forming regions hosting the masers, and we find that these regions lag circular rotation on average by $\\sim$17~km~s$^{-1}$, a value comparable to those found recently by similar studies. } {}  \\keywords {Techniques: interferometric -- Astrometry -- Masers -- Stars: formation -- ISM: molecules -- Galaxy: kinematics and dynamics} ",
        "introduction": "Accurate distances and proper motions are crucial for studies of the structure and kinematics of the Milky Way. Since we are in the Galactic plane, it is not easy to determine the spiral structure of our Milky Way \\citep{reid:2009b}. Massive star-forming regions trace the spiral arms and are objects well-suited to revealing the structure of the Milky Way.  Determining the fundamental physical properties of individual objects, such as size scales, masses, luminosities, and ages, also depends critically on distance. For example, the distance to the Orion Nebula, determined by a number of trigonometric parallax measurements using radio continuum emission from stars in its associated cluster \\citep{menten:2007,Sandstrom:2007}, water masers \\citep{hirota:2007}, and SiO masers \\citep{kim:2008}, turned out to be 10\\% less than previously assumed, resulting in 10\\% lower masses, 20\\% fainter luminosities, and 20--30\\% younger ages for the stars in the cluster.   The most fundamental and unbiased method of measuring distances is the trigonometric parallax, which depends only on geometry and is therefore free of any astrophysical assumptions. To achieve the sub-milliarcsecond astrometric accuracy at optical wavelengths requires space-borne observations. The first dedicated optical satellite for this purpose was ESA's Hipparcos \\citep{perryman:1995}. With its accuracy of 0.8--2\\,mas, it was capable of measuring distances up to $\\sim$200\\,pc -- only a small fraction of the Milky Way. A new optical astrometry satellite GAIA, to be launched in 2012, will be two orders of magnitude more accurate and push the optical parallaxes to Galactic size scales \\citep[e.g.,][]{lindgren:2009}. Despite the high accuracy, GAIA will suffer from dust extinction in the Galactic plane, in particular in the spiral arms and toward the Galactic center. Here radio astronomy can provide an important complement to GAIA, because radio wavelengths do not suffer from dust extinction, and Very-Long-Baseline Interferometry (VLBI) phase-referencing techniques can provide accuracies better than 10\\,$\\mathrm{\\mu}$as, allowing accurate distances with errors less than 10\\% out to 10\\,kpc {\\cite[see for example][]{honma:2007,hachisuka:2009,reid:2009a}.  Strong and compact radio sources, such as molecular masers, are ideal targets for trigonometric parallax measurements. Masers are frequently found in (dusty) star formation regions (SFRs) and asymptotic giant branch stars. The most numerous interstellar masers are the 22\\,GHz water masers, followed in number by the 6.7\\,GHz methanol masers; to date several hundreds of these masers have been found across the Galaxy \\citep[e.g.,][]{pestalozzi:2005, green:2009}. They are exclusively associated with early stages of massive star formation, observed both prior to the development of an ultra compact H{\\sc ii} (UCH{\\sc ii}) region and coexistent with these \\citep{menten:1991, ellingsen:2006, pestalozzi:2007}. Maser modeling indicates that the 6.7\\,GHz emission is likely to be excited via radiative pumping by warm dust heated by newly formed high-mass stars \\citep{sobolev:2007}.  Sources of methanol masers have long lifetimes of $\\sim$\\,$10^4$\\,year \\citep{walt:2005}. Their velocity spread is typically within $\\pm5$\\,km~s$^{-1}$ about the systematic velocity of the molecular core, and single maser features are usually narrow ($\\sim$1\\,km~s$^{-1}$). The kinematics suggest that methanol masers originate from different regions than water masers, which have wider velocity spreads and are produced in protostellar outflows \\citep{menten:1996}. To summarize, the strength, ubiquity, long lifetimes, and small internal motions make the 6.7\\,GHz methanol maser very suitable for astrometric purposes.  Here we present results from observations with the European VLBI Network (EVN) of eight strong methanol masers belonging to well-known massive SFRs in the outer part of the Milky Way: ON\\,1, L\\,1287, L\\,1206, NGC\\,281-W, MonR\\,2, S\\,252, S\\,255 and S\\,269. Recently, H$_2$O maser or 12.2\\,GHz methanol maser parallax measurements have been reported for three of these regions (NGC\\,281-W, \\citeauthor{sato:2008} \\citeyear{sato:2008}; S\\,252, \\citeauthor{reid:2009a} \\citeyear{reid:2009a}; S\\,269, \\citeauthor{honma:2007} \\citeyear{honma:2007}), allowing a valuable cross check with our measurements. The results of this present work are crucial for star formation studies in these regions, and, together with a larger sample of parallaxes, will help to understand the structure of the Local and Perseus arms.  In this paper we report on the first parallax measurements made with the EVN. Results presented here replace preliminary, less accurate, and less complete measurements \\citep{rygl:2008}. In the next section we describe the observations and data analysis, and in Sect.~3 we explain the method and fitting of the parallax and proper motion. The results are then presented per source in Sect.~4 and discussed in Sect.~5. The conclusions are summarized in Sect.~6. ",
        "conclusions": "\\subsection{Space motions}  \\begin{table*}[!htb] \\caption{\\label{ta:3d}Peculiar motions} \\centering \\begin{tabular}{lr@{$\\pm$}lr@{$\\pm$}lr@{$\\pm$}lr@{$\\pm$}lr@{$\\pm$}lr@{$\\pm$}l} \\hline\\hline Source & \\multicolumn{6}{c}{$R_0=8.5$ kpc, $\\Theta_0=220~\\mathrm{km~s^{-1}}$} & \\multicolumn{6}{c}{$R_0=8.4$ kpc, $\\Theta_0=254~\\mathrm{km~s^{-1}}$}\\\\ & \\multicolumn{2}{c}{$U$} & \\multicolumn{2}{c}{$V$} & \\multicolumn{2}{c}{$W$} & \\multicolumn{2}{c}{$U$} & \\multicolumn{2}{c}{$V$} & \\multicolumn{2}{c}{$W$}\\\\ & \\multicolumn{2}{l}{($\\mathrm{km~s^{-1}}$)}& \\multicolumn{2}{l}{($\\mathrm{km~s^{-1}}$)}&\\multicolumn{2}{l}{($\\mathrm{km~s^{-1}}$)} & \\multicolumn{2}{c}{($\\mathrm{km~s^{-1}}$)}& \\multicolumn{2}{c}{($\\mathrm{km~s^{-1}}$)}&\\multicolumn{2}{c}{($\\mathrm{km~s^{-1}}$)}\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} ON\\,1    & $18$&$8$ & $-19$&3 & $4$& $10$ &$7$&$8$ & $-21$&$3$& $4$&$10$ \\\\ L\\,1206   & $2$&$4$ & $-16$&$2$ & $0$&$6$  &$-1$&$4$ & $-16$&$2$& $0$&$6$\\\\ L\\,1287   & $13$&$2$ & $-18$&$2$ & $-3$&$2$  &$10$&$2$ & $-18$&$2$ &$-3$&$3$\\\\ NGC\\,281-W& $13$&$2$ & $-3$&$2$ & $-9$&$2$  &$6$&$2$& $-3$&$2$ &$-9$&$2$\\\\ S\\,255    & $1$&$3$ & $-4$&$12$ & $3$&$7$  & $2$&$3$ & $-4$&$12$ & $3$&$7$\\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{table*}  \\begin{figure} \\centering \\includegraphics[width=\\columnwidth]{1313520.eps} \\caption{\\label{fig:gal}An artist's impression of a plane-on view of our Galaxy (image credit: R. Hurt NASA/JPL-Caltech/SSC), overlaid with parallax measurements of water and methanol masers between $60\\degr<l<240\\degr$ taken from \\citet{reid:2009b}, in which recent parallax measurements have been put together to study the Galactic structure. The black dot marks the Galactic center, the yellow dot encircled with black the Sun, the yellow dots mark the parallaxes taken from the literature, the orange dots encircled with white, and white arrows mark the parallaxes and peculiar motions obtained in this work. The peculiar motions are shown after removing the Galactic rotation assuming the values of $R_0=8.4$ kpc and $\\Theta_0=254~\\mathrm{km~s^{-1}}$.} \\end{figure}  With parallax and proper motion measurements, one can calculate the full space motion of the masers in an SFR in the Galaxy.  Using an accurate model of Galactic dynamics (Galactic rotation speed, ${\\Theta_0}$, distance of the Sun to the Galactic center, ${R_0}$, rotation curve) for a source and removing the modeled contribution of Galactic rotation allows one to retrieve source peculiar motion relative to a circular orbit. The peculiar motion is described by the vectors $U$, $V$, and $W$, locally toward the Galactic center, toward the direction of rotation and toward the North Galactic Pole, respectively.  In Table \\ref{ta:3d}, the peculiar motions of our sources are given for two different Galactic models: the I.A.U.~recommended values for the LSR motion $R_0=8.5$ kpc, $\\Theta_0=220~\\mathrm{km~s^{-1}}$ in columns 2--4; a new model with $R_0=8.4$ kpc, $\\Theta_0=254~\\mathrm{km~s^{-1}}$, based on a large sample of trigonometric parallax measurements \\citep{reid:2009b}, in columns 5--7. Both models adopt the Hipparcos values for the Solar motion \\citep{dehnen:1998} and assume a flat rotation curve. We used the parallaxes and proper motions as are listed in Table \\ref{ta:sum}, where the uncertainty introduced by the apparent movement of the background sources was included in the proper motion uncertainty of the source.  In Fig.~\\ref{fig:gal} we plot the SFRs studied in this work at their determined distances with their peculiar motion in the Galactic plane, after removing the Galactic rotation. For either model, it is clear that three SFRs show a similar lag on circular rotation as the SFRs studied by \\citet{reid:2009b}, while two have a circular rotation close to zero lag. For NGC\\,281-W, the zero lag can be understood from its large distance from the Galactic disk ($b\\simeq-6\\degr$) and the dominant contribution from the expanding super bubble to the peculiar motion of this SFR (see Sect.~\\ref{sec:ngc}). For S\\,255, it is not clear why its circular velocity is close to Galactic rotation (Sect.~\\ref{sec:s255}). The motion towards the Galactic center, $U$, tends to be positive for most of the SFRs.  Recently, \\citet{mcmillan:2009} have published a reanalysis of the maser astrometry presented by \\citet{reid:2009b}. They fitted the data with the revised solar peculiar motion of Binney~(2009), mentioned in \\citet{mcmillan:2009}, where the $V$ component of the solar peculiar velocity, $V_\\odot$, is increased from 5.2~km~s$^{-1}$ to 11~km~s$^{-1}$. With a larger $V_\\odot$ the lag of SFRs on Galactic rotation would decrease to $\\sim11~\\mathrm{km~s^{-1}}$, which agrees with the expected velocity dispersion for young stars of $\\sim10~\\mathrm{km~s^{-1}}$ \\citep{aumer:2009}. However, even when considering the revised $V_\\odot$, we note that most SFRs still rotate more slowly than the Galactic rotation (by $\\sim11~\\mathrm{km~s^{-1}}$) and that their velocities are not randomly dispersed.   \\subsection{Onsala 1} Our trigonometric parallax measurement places ON\\,1 on a distance from the Sun of $2.57^{+0.34}_{-0.27}$\\,kpc, somewhat closer than the (near) kinematic distance of 3.0\\,kpc, based on the methanol maser line at $15\\,\\mathrm{km~s^{-1}}$. The latter assumes a flat rotation curve with as distance of the Sun to the Galactic center, $R_0=8.5$\\,kpc, and for the Galactic rotation speed, $\\Theta_0=220$\\,km~s$^{-1}$. However, it is farther than the commonly adopted (near) kinematic distance of 1.8\\,kpc based on formaldehyde at $v_\\mathrm{LSR}=11.2\\,\\mathrm{km~s^{-1}}$ \\citep{macleod:1998} using the Galactic rotation curve of \\citet{wouterloot:1989}, with  $R_0=8.5$\\,kpc, $\\Theta_0=220$~km~s$^{-1}$. Our result resolves the near/far kinematic distance ambiguity and locates ON\\,1 in the Local spur consistent with the near kinematic distance.  We found different proper motions for the northern and southern maser groups (Table \\ref{ta:data}). The resulting difference in the east-west direction was small, $\\mu_\\alpha=0.18\\pm0.24\\,\\mathrm{mas~yr^{-1}}$, but in the north-south direction, the southern group was moving by $\\mu_\\delta=-0.77\\pm0.12\\,\\mathrm{mas~yr^{-1}}$ away from the northern group. At a distance of 2.6\\,kpc, this corresponds to a relative speed of 9.4\\,km~s$^{-1}$.  This large proper motion may be explained by the masers being located in the molecular gas surrounding an expanding H{\\sc ii} region, as suggested by \\citet{fish:2007} and \\citet{su:2009}. The H{\\sc ii} region is located at ($\\alpha=20^{\\mathrm{h}}10^{\\mathrm{m}}09\\rlap{$.$}\\,^{\\mathrm{s}}03,~\\delta= +31\\degr 31\\arcmin 35\\rlap{$.$}\\,\\arcsec 4$, J2000) and between the two maser groups. The radial velocity of the H{\\sc ii} region, from the H76$\\alpha$ recombination line, is $5.1 \\pm 2.5$~km~s$^{-1}$ \\citep{zheng:1985}. The methanol masers would be (with respect to the rest frame of ON\\,1 at $5.1$\\,km~s$^{-1}$) in a blue-shifted component at $\\sim$0 km~s$^{-1}$, northward, and a red-shifted component at $\\sim$15 km~s$^{-1}$, southward of the H{\\sc ii} region. Also, the hydroxyl masers in ON\\,1 provide some confirmation of this scenario \\citep{namm:2006,fish:2007}. The expansion velocities of the masers, calculated by assuming a linear expansion from the center of the H{\\sc ii} region, would be 5.8\\,km~s$^{-1}$ and 3.6\\,km~s$^{-1}$ for the northern and southern maser groups, respectively.  \\subsection{L\\,1206}  For L\\,1206, the trigonometric parallax distance, $0.776^{+0.104}_{-0.083}$~kpc, is shorter than the kinematic distance; 1.0 kpc \\citep[based on a hydroxyl maser, at $v_\\mathrm{LSR}=-8.5~\\mathrm{km~s^{-1}}$, using the Galactic rotation curve of \\citet{wouterloot:1989} with $R_0=8.5$~kpc, $\\Theta_0=220$~km~s$^{-1}$,][]{macleod:1998}, 1.23 kpc \\citep[ammonia, $v_\\mathrm{LSR}=-9.9~\\mathrm{km~s^{-1}}$, using the Galactic rotation curve of \\citet{brand:1993}, with $R_0=8.5$~kpc, $\\Theta_0=220$~km~s$^{-1}$][]{molinari:1996} and 1.4 kpc (6.7 GHz methanol, flat rotation curve, $R_0=8.5$~kpc, $\\Theta_0=220$~km~s$^{-1}$). Our result places L\\,1206 in the Local spur. Cep A, a nearby ($3\\rlap{$.$}\\,\\degr8$\\ separation, P.A. 126 \\degr) star-forming region, is located at a distance of $700\\pm39$ pc, as determined by a parallax measurement of 12.2 GHz methanol masers \\citep{moscadelli:2009}. The L\\,1206 SFR therefore seems to be in the same part of the Local arm as Cep A. L\\,1206 is a dark cloud with an infrared source, IRAS 22272+6358A, which coincides with the methanol maser emission and a 2.7\\,mm dust continuum peak \\citep{beltran:2006}. Since neither 2\\,cm nor 6\\,cm radio emission is detected \\citep{wilking:1989,McCutcheon:1991}, L\\,1206 is thought to be in a young phase prior to the formation of an H{\\sc ii} region. \\citet{beltran:2006} report large CO outflows, and put the systemic velocity of the ambient medium between [$-13.5$, $-8.5$] km s$^{-1}$, consistent with our systemic velocity range of [$-13.3$, $-10.9$] km s$^{-1}$.  \\subsection{L\\,1287}  The parallax for L\\,1287 sets it at a distance of $0.929^{+0.034}_{-0.033}$\\,kpc. This is close to the photometric distance of $\\sim$850\\,pc from \\citep{yang:1991}, which placed L\\,1287 in the Local arm. However, the kinematic distance based on the methanol maser line would place L\\,1287 at 2\\,kpc in the Perseus arm.  Methanol masers in the dark cloud L\\,1287 are located at the base of the bipolar CO outflow \\citep{yang:1991} originating in the main core of the cloud. The infrared point source (IRAS 00338+6312) in the center of the core indicates a proto-stellar object surrounded by cold, high density  $\\mathrm{NH_3}$ (1,1) gas \\citep{estalella:1993}.  \\subsection{NGC\\,281-W} \\begin{figure} \\centering \\includegraphics[width=10cm,angle=-90]{1313521.eps} \\caption{Lower panel: a face-on view of the Galactic plane. The Galactic longitudes of NGC\\,281-W ($l\\simeq123\\degr$) and IRAS 00425530 ($l\\simeq122\\degr$) are indicated by dashed lines. The top panel shows a cross-section of the Galactic plane along the Galactic longitude of NGC\\,281-W (123\\degr). The vertical coordinate of this plot is the direction perpendicular to the Galactic plane. We plot NGC\\,281-W and IRAS 00420+5530 as black dots, together with the respective distance error. Black arrows show the peculiar motion of the source after removing the Galactic rotation ($R_0=8.4$ kpc, $\\Theta_0=254~\\mathrm{km~s^{-1}}$). The errors in peculiar motion are indicated by the gray cones. The light-gray circles mark the super bubble with the center of the bubble at a distance of 2.5 kpc and a radius of 0.4 kpc, adopted after \\citet{sato:2008}. The cross marks the averaged distance of 2.58~kpc. The dotted lines show the backward prolongation of the proper motion of NGC\\,281-W and IRAS~00420+5530.} \\label{fig:bub} \\end{figure}  \\label{sec:ngc} There are numerous distance estimates for NGC\\,281-W in the literature. Optical photometry places the cloud between 2 and 3.68 kpc \\citep{sharpless:1954, cruz-g:1974,  henning:1994, guetter:1997}, while the kinematic distance is 3 kpc \\citep{lee:2003}, based on a $v_\\mathrm{LSR}=-30\\,\\mathrm{km~s^{-1}}$, using the Galactic rotation curve of \\citet{clemens:1985} and $R_0=8.5\\,\\mathrm{kpc},\\theta_0=220\\,\\mathrm{km~s^{-1}}$. The radial velocity of the methanol maser line at $-29~\\mathrm{km~s^{-1}}$, assuming a flat Galactic rotation curve, and $R_0=8.5\\,\\mathrm{kpc},\\theta_0=220\\,\\mathrm{km~s^{-1}}$, would place NGC\\,281-W at 2.5 kpc.  Recently, \\citet{sato:2008} reported a water maser parallax value of $0.355\\pm0.030$\\,mas, measured with the VERA interferometer, corresponding to a distance of $2.82\\pm0.24$\\,kpc. We measured a parallax of $0.421\\pm0.022$\\,mas, arriving at a distance of $2.38^{+0.13}_{-0.12}$\\,kpc. These distance measurements agree with \\citet{sato:2008} within 2$\\sigma$ of the joint uncertainty.  Also, our proper motion agrees within 1$\\sigma$ with the proper motion by \\citet{sato:2008}, and our radial velocity ($v_\\mathrm{LSR}\\sim-29$\\,km~s$^{-1}$) is close to that of the water masers ($v_\\mathrm{LSR}\\sim-31$\\,km~s$^{-1}$). It is unlikely that the water and methanol masers originate at such different depths into the SFR that it would affect the distance. For example, in the SFR W3(OH), the water and methanol maser parallaxes are in good agreement \\citep{xu:2006,hachi:2006}.  NGC\\,281-W is an SFR that lies near the edge of a super bubble in the Perseus arm \\citep{megeath:2003,Sato:2007}. Our slightly shorter distance would change the position of NGC\\,281-W on this super bubble. Another SFR located on this bubble has a parallax determined distance: IRAS 00420+5530 at $2.17\\pm0.05$~kpc \\citep{moellen:2009}. This allows to compare the distances and peculiar motions of both SFRs, and check whether the expanding super bubble, with the expansion center at 2.5~kpc adopted after \\citet{sato:2008}, can be responsible for the peculiar motions of both of them. We calculated the peculiar motions of IRAS 0042+5530 after \\citet{moellen:2009} and  NGC 281-W after \\citet{sato:2008} and our results. The bottom panel of Fig.~\\ref{fig:bub} shows the peculiar motions in a face-on view of the Galactic plane. From the figure can be seen that the $U$ component of NGC\\,281-W, representing the vector toward the Galactic center, is positive and non-zero in both our study and the work of \\citet{sato:2008}. This means that NGC\\,281-W is moving toward the Galactic center, which is the expected movement for the near side of an expanding bubble. In the top panel of Fig.~\\ref{fig:bub} the peculiar motions of both SFRs are shown projected on a cross-section of the Galactic plane along the longitude of NGC 281-W, $l=123\\degr$. The average distance to NGC\\,281-W based on \\citet{sato:2008} and this work, $2.58^{+0.27}_{-0.23}$~kpc, is marked in this panel as a cross. If we prolong the peculiar motion backwards of both IRAS 00420+5530 and the average result for NGC 281-W, they intersect at a distance from the Sun of 2.65~kpc and at a Galactic latitude of $\\sim$$-0\\rlap{$.$}\\,\\degr1$. This suggets that, if indeed both SFRs are expanding linearly from one expansion center, the expansion center is offset from the previously assumed position. However, this offset center of expansion is difficult to understand when looking at the peculiar motion of both SFRs in the Galactic plane (Fig.~\\ref{fig:bub}, bottom panel), since the motion does not seem to originate in one mutual expansion center; hence, it is difficult to pinpoint the characteristics of the super bubble if the peculiar motions and distances of only two SFRs are all that are known.     \\subsection{S\\,255} \\label{sec:s255} For S\\,255, the parallax distance, $1.59^{+0.07}_{-0.06}$\\,kpc, is much closer than the commonly used photometric distance of 2.5\\,kpc \\citep{moffat:1979,blitz:1982}. S\\,255 is an individual H{\\sc ii} region, associated with a complex of H{\\sc ii} regions. The methanol maser emission coincides with a filament of cold dust and molecular gas between two H{\\sc ii} regions, S\\,255 and S\\,257. \\cite{minier:2007} studied the star formation in this filament, which they propose is possibly triggered by the compression of the filament by the two H{\\sc ii} regions. They find several molecular clumps in this filament. At distance of $\\approx2.5$\\,kpc, the masses of the clumps, determined from the submillimeter dust continuum, are around 300\\,$\\mathrm{M_\\odot}$ \\citep{minier:2007}. However, if we place S\\,255 at the parallax determined distance of 1.6\\,kpc, the clump masses would drop by 60\\% to $\\sim120$\\,$\\mathrm{M_\\odot}$."
    },
    "0910/0910.3918_arXiv.txt": {
        "abstract": "In this paper we consider a simple two-fluid model for pulsar glitches.  We derive the basic equations that govern the spin evolution of the system from two-fluid hydrodynamics, accounting for the vortex mediated mutual friction force that determines the glitch rise. This leads to a simple ``bulk'' model that can be used to describe the main properties of a glitch event resulting from vortex unpinning. In order to model the long term relaxation following the glitch our model would require additional assumptions regarding the repinning of vortices, an issue that we only touch upon briefly. Instead, we focus on comparing the phenomenological model to results obtained from time-evolutions of the linearised two-fluid equations, i.e. a ``hydrodynamic'' model for glitches. This allows us to study, for the first time, dynamics that was ``averaged'' in the bulk model, i.e. consider the various neutron star oscillation modes that are excited during a glitch.  The hydro-results are of some relevance for efforts to detect gravitational waves from glitching pulsars, although the conclusions drawn from our rather simple model are pessimistic as far as the detectability of these events is concerned. ",
        "introduction": "Neutron stars provide excellent testbeds for extreme physics theory. Since their core density reaches beyond anything we can produce in the laboratory, observations of such compact remnants may provide unique constraints on models for supranuclear physics \\citep{LP04}. In order to infer useful information from astrophysical observations we need to construct realistic models with the power to predict the evolution of individual systems (at least to some extent).  This is a serious challenge for the modelling community.  Even a moderately reasonable neutron star model should account for the presence of different exotic states of matter.  From nuclear physics and Bardeen-Cooper-Schrieffer (BCS) theory we expect the outer neutron star core to consist of superfluid neutrons, superconducting protons, free electrons and muons. Deeper into the core more exotic phases of matter, like superfluid hyperons and/or deconfined quarks exhibiting colour-flavour-locked superconductivity, are likely to be present. Meanwhile, a relatively thin (1 km or so) crust surrounds the fluid core.  In the crust, matter changes from a soup of nucleons in the interior to an elastic nuclear lattice of heavy iron near the surface. In the inner part of the crust (beyond neutron drip) free neutrons are expected to be superfluid.  An important question concerns whether differences in internal structure can be deduced from observations. This is obviously problematic since most of the data collected for these stars originate from the star's surface, atmosphere or magnetosphere. Having said that, there are phenomena that involve bulk dynamics and which should depend on the internal composition. Examples of such phenomena are the radio pulsar glitches \\citep{2000MNRAS.315..534L} and the quasi-periodic oscillations observed in the X-rays from magnetar flares, see for example \\citet{2007Ap&SS.308..625W}. The long relaxation time associated with glitches is seen as indirect evidence for neutron star superfluidity \\citep{AI75,Rud76,Alp81}, and recent considerations of the crust oscillations which are thought to be the origin of the observed magnetar oscillations suggest that these may be affected by the presence of a crust superfluid as well \\citep{AGS09}.  In this paper we discuss basic models for pulsar glitches. We focus on the glitch event itself rather than the subsequent long term relaxation. We consider the implications of the standard two-fluid model from two different points of view. First of all, we derive the simple equations that govern the bulk evolution of the system from two-fluid hydrodynamics, accounting for the vortex mediated mutual friction force. This leads to a basic model that can be used to describe a glitch rise resulting from global vortex unpinning, be it in the crust \\citep{1996ApJ...457..844L,2008ApJ...672.1103M,2008MNRAS.390..175W,2009ApJ...700.1524M} or in the core \\citep{2003PhRvL..91j1101L}\\footnote{In reality, our model is somewhat unrealistic for both crust and core superfluids. In the former case we have not accounted for the crust elasticity, while in the latter case we are ignoring the expected interaction between neutron vortices and proton fluxtubes. These effects may have a decisive impact on glitch dynamics. Nevertheless, the present work is state-of-the-art in this area, and we expect to add the relevant features to the hydrodynamical model in future work. }. The results demonstrate that the model requires additional assumptions regarding the repinning of vortices in order to model the long-term evolution. This is as expected \\citep{1984ApJ...276..325A}. Secondly, we use time-evolution of the linearised two-fluid equations as a ``hydrodynamic'' model for the glitch event.  This allows us to consider dynamics that was ``averaged'' in the bulk model. In particular, we consider the various neutron star oscillation modes that are excited during the glitch. In principle, the obtained results could be of relevance for efforts to detect gravitational waves from glitching pulsars. Having said that, our estimates suggest that the gravitational-wave signals associated with the impulsive events that we consider here are very weak. ",
        "conclusions": "We have discussed the dynamics of pulsar glitch events from two, complementary, points of view. First we constructed a simple model based on global ``averaging'' of the standard two-fluid equations including the mutual friction due to superfluid vortices. This analysis provides a more detailed derivation of the phenomenological relations that have been used in many discussions of glitches.  In particular, our final relations clarify how the spin-up time depends on key parameters like the entrainment. The derivation also highlights the various assumptions and the restricted validity of the model. Anyway, for typical values of the parameters (see Sec.~\\ref{sec:4.3}), our model has a glitch rise time shorter than the upper bound set by current observations.  The model provides a useful description of the actual glitch event, but it does not account for the subsequent long-term relaxation (on a timescale of days to months) of the system.  A key conclusion from our discussion is that the late stages of evolution requires additional assumptions, most likely, concerning the repinning of vortices. Understanding this phase better, e.g. connecting it to the two-fluid hydrodynamics and the averaged forces that act on the vortices, is an important challenge for the future.  It seems clear that vortex creep will play a central role \\citep{AI75,1984ApJ...276..325A,1989ApJ...346..823A,1993ApJ...403..285L,1993ApJ...409..345A}, but this mechanism has not yet been discussed in terms of the macroscopic hydrodynamics.  This issue needs to be addressed if we are to develop more detailed models of glitch dynamics. We definitely need to move beyond phenomenology.  As a first step towards hydrodynamic glitch modelling, we have extended the recent linear perturbation evolution code of~\\citet{2009MNRAS.396..951P} to include the mutual friction and the perturbed gravitational potential.  Initiated with perturbations that represent two fluids rotating uniformly at different rates, the numerical code shows how the system relaxes to co-rotation. We have analysed this relaxation in detail and demonstrated that the behaviour is accurately described by the phenomenological model, at least for non-stratified stellar models. When the star is stratified (e.g. has varying composition) the relaxation deviates from the simple model. This is as expected, since the global model was derived under the assumption of uniform parameters. Of course, the numerical evolutions provide us with a useful tool for studying the behaviour of more complex stellar models. In addition, our time evolutions provide a first insight into the excitation of neutron star oscillations by glitches. Our results show that a set of axisymmetric modes are excited by the glitch initial data.  These modes will radiate gravitational waves\\footnote{In principle, the induced oscillations may also lead to variations in the electromagnetic signal.  However, in order to quantify this effect one would need a more detailed analysis of the coupling between the motion of the crust and the magnetosphere. Such estimates are beyond the scope of the present analysis, but it is worth noting that the oscillation modes that we consider are all in the kHz range (much faster than the spin-period) meaning that they would not be seen as ``modulations'' of the pulsar signal. }, and it is important to establish if the associated signals may be observable with future detectors. In this respect, our results are quite pessimistic.  In the cases that we have considered, the gravitational-wave signal is too weak to be detectable (even with a third generation of detectors)\\footnote{At first sight this conclusion seems at variance with the results of \\citet{2008CQGra..25v5020V}, who consider a different glitch scenario. In their (cylindrical) model problem the gravitational-waves are associated with the large scale Ekman flow that results from a rotational lag between the crust and the core in the star. The two mechanisms are obviously different. In particular, in our case the event is impulsive and one would expect the signal to be burst-like. We certainly cannot envisage the 14 day integration suggested by \\citet{2008CQGra..25v5020V} to improve the signal to noise ratio.  Basically, the model parameters used by \\citet{2008CQGra..25v5020V} seem rather optimistic.}. However, it is not clear that this is the final say on the matter. One should keep in mind that the gravitational-wave strain differs enormously for our two model configurations. The non-stratified model does not (in principle) radiate at all, while the stratified model leads to a qualitatively interesting (albeit weak) signal.  The enormous difference between these results shows that we need to continue to refine our modelling. We obviously have to account for the variation of composition throughout the star, and consider the fact that superfluid components will only be present in specific density regions. We also need to understand the nature of the vortex pinning better. In our models we have assumed that the vortices unpin in a catastrophic global event. It is far from clear that this is the case in a real system. It could, for example, be that the unpinning is localized. This would make the event less symmetric which may enhance the gravitational-wave signal. We clearly need to understand the actual mechanism that triggers the glitches better.  The superfluid instability discussed by \\citet{2009PhRvL.102n1101G} is interesting in this respect, but we need to study this mechanism in more detail to establish to what extent it can operate in a real neutron star.  To make progress we need to overcome a number of challenges. Yet, recent developments have provided us with interesting insights and (most importantly) computational technology that should allow us to study much more realistic neutron star models in the not too distant future."
    },
    "0910/0910.5225_arXiv.txt": {
        "abstract": "V371~Per was found to be a double-mode Cepheid with a fundamental mode period of 1.738 days, the shortest among Galactic beat Cepheids, and an unusually high period ratio of 0.731, while the other Galactic beat Cepheids have period ratios between 0.697 and 0.713. The latter suggests that the star has a metallicity [Fe/H] between -1 and -0.7. The derived distance from the Galactic Plane places it in the Thick Disk or the Halo, while all other Galactic beat Cepheids belong to the Thin Disk. There are indications from historical data that both the fundamental and first overtone periods have lengthened. ",
        "introduction": "The variability of V371~Per was discovered by \\citet{weber}. From photographic studies, \\citet{satyvaldiev} and \\citet{meinunger} did not find any periodicity, and the star was therefore assumed to be an irregular variable. \\citet{behlen} found a period of 1.2697 days from CCD data and suspected it to be a beat Cepheid because of its changing light curve. This interpretation agrees with the spectral type G0 given by \\citet{bond}. As a consequence, the star was observed for a number of years at the Behlen Observatory at the University of Nebraska, the United States Naval Observatory Flagstaff Station (NOFS), the Sonoita Research Observatory (SRO; Sonoita, Arizona) and the Zagori Observatory (Athens, Greece).  The resulting data will be presented in this paper. In addition to the new data, the data sets from \\citet{satyvaldiev} and \\citet{behlen} are reanalysed (the latter with revised values for the comparison stars), and data from the Northern Sky Variability Survey \\citep[NSVS,][]{nsvs} are analysed as well. ",
        "conclusions": "Extensive CCD photometry of V371~Per over a number of years has clearly shown it to be a Galactic beat Cepheid, with the shortest period known so far. The high value of the frequency ratio (0.731) suggests that its metallicity is much lower than that of the other Galactic beat Cepheids. Its distance derived from empirical PL relations places it in the Galactic Thick Disk or Halo, 0.8~kpc above the Galactic Plane. V371~Per is therefore a remarkable object and a spectroscopic study is warranted to confirm its low metallicity. It is very likely that V371~Per will help to refine the theoretical models for Cepheid pulsation. Further photometric observations may also confirm the period increase suspected for both modes. This will become evident within the next five to ten years."
    },
    "0910/0910.5425.txt": {
        "abstract": "{} {This work reports radiative transition rates and electron impact excitation collision strengths for levels of the 3s$^2$3p, 3s3p$^2$, 3s$^2$4s, and 3s$^2$3d configurations of \\ion{Si}{ii}.}{The radiative data were  computed using the Thomas-Fermi-Dirac-Amaldi central potential, but with the modifications introduced by Bautista (2008) that account for the effects of electron-electron interactions. We also introduce new schemes for the optimization of the variational parameters of the potential. Additional calculations were carried out with the Relativistic Hartree-Fock and the multiconfiguration Dirac-Fock methods. Collision strengths in LS-coupling were calculated in the close coupling approximation with the R-matrix method. Then, fine structure collision strengths were obtained by means of the intermediate-coupling frame transformation (ICFT) method which accounts for spin-orbit coupling effects. }{ We present extensive comparisons between the results of different approximations and with the most recent calculations and experiment available in the literature. From these comparisons we derive a recommended set of $gf$-values and radiative transition rates with their corresponding estimated uncertainties. We also study the effects of different approximations in the representation of the target ion on the electron-impact collision strengths. Our most accurate set of collision strengths were integrated over a Maxwellian distribution of electron energies and the resulting effective collision strengths are given for a wide range of temperatures. Our results present significant differences from recent calculations with the B-spline non-orthogonal R-matrix method. We discuss the sources of the differences. } {} ",
        "introduction": "Singly ionized silicon (\\ion{Si}{ii}) is prominent in ultra-violet (UV) and optical spectra of various astrophysical plasmas. In terms of absorption spectra, in the spectral range longward of 912 \\AA\\ \\ion{Si}{ii} has 8 absorption lines complexes connected to the ground term 3s$^2$3p $^2$P$^o$. In photoionized plasmas with electron temperatures ($T_e$) of the order of $10^4$ K and electron densities ($n_e$) $\\le 10^4$ cm$^{-3}$ the relative optical depths of troughs from the two ground state levels depend on $n_e$. Hence, this dependence can be used as diagnostic of $n_e$ (e.g. Dunn et al. 2009). On this regard, the lines centered at 1814 \\AA (3s$^2$3p $^2$P$^o$ - 3s3p$^2$ $^2$D) are particularly convenient because of their unusually small oscillator strengths, such that these, among all other \\ion{Si}{ii} lines, are the most likely to be in the linear part of the curve of growth and their column densities can be accurately measured. Unfortunately, though, the determination of accurate and reliable oscillator strengths for these transitions is particularly difficult and has been the subject of much theoretical and experimental efforts.  \\ion{Si}{ii} is also prominent in emission spectra of various kinds of objects. In the upper chromosphere and lower transition region in the Sun and late-type stars, line ratios among the \\ion{Si}{ii} 1814 \\AA\\ multiplet and the intercombination (3s$^2$3p $^2$P$^o$ - 3s3p$^2$ $^4$P) multiplet near 2335 \\AA\\ are potentially useful density diagnostics. However, until recently the best electron impact collision strengths of \\cite{dufton} led to predicted line ratios that disagreed with observations \\citep{judge}. Similarly, theoretical models of emission spectra of Broad Line Regions of Active Galactic Nuclei fail to reproduce the observed intensities by factors of a few; a problem that was named by \\cite{baldwin} the ''\\ion{Si}{ii} disaster''.  Despite considerable theoretical work on oscillator strengths there is still considerable spread in the results, particularly for those transitions of most astronomical interest, such as those of the 1814 \\AA\\ complex. This is because the upper $^2$D term of these transitions is made of a mixture of the 3s3p$^2$ and 3s$^2$3d configurations, which produce strong cancellation in the oscillator strengths and makes the f-values very difficult to compute. Also for these transitions, there has been much spread in experimental gf-values. Measurements based on the electron beam phase shift method \\citep{savage, curtis} could be inaccurate for very long lifetimes. Absolute emission measurements on an arc \\citep{hofmann} are often uncertain owing to the difficult control and calibration of the instrument. Determinations based in comparisons of equivalent widths of the weak 1810 \\AA\\ lines to  stronger lines in astronomical spectra \\citep{shull, vanburen} are unreliable due to blends and saturation of the stronger lines \\citep{jenkins}. The latest, and probably most accurate determination of gf-values for these lines was done with the time-resolved laser-induced fluorescence technique \\cite{bergeson}.  The oscillator strengths for the intercombination transitions (3s$^2$3p $^2$P$^o$ - 3s3p$^2$ $^4$P) have also been subject of controversy. These transitions arise due to spin-orbit mixing of the metastable term 3s3p$^2$ $^4$P with the $^2$S, $^2$P, and $^2$D terms. The difficulty then comes from the fact that the gf-values are sensitive to core-valence and core-core correlations, in addition to the valence-shell effects. The results of some of the most recent calculations \\citep{froesefischer} differ from experimentally measured transition probabilities \\citep{calamai} by $\\sim$ 30\\%.  But beyond providing accurate oscillator strengths by tailoring atomic structure representations on each type of transition of interest, it is important to try constructing a single representation that yields reasonably accurate energies and oscillator strengths for all levels and transitions considered simultaneously. This is because various practical applications will require such a general atomic representation for subsequent scattering calculations (e.g. electron impact excitation, photoionization, and recombination). One of the most widely used methods for this purpose is based on the Thomas-Fermi-Dirac-Amaldi (TFDA) central potential to generate optimized one-electron orbitals, which represent the atomic structure through configuration interaction (CI) expansions and can also be used for close-coupling representations of the scattering problem. Recently, \\cite{bautista08} introduced a correction to the TFDA potential that accounts for the effects of electron-electron interactions on the radial wavefunction and yield a considerably improved representation of the system. It is thus interesting to apply this new approach to get the best possible representation of the important \\ion{Si}{ii} system.  In the present paper, to complement and expand the work with the TFDA potential we also compare with the results of the Hartree-Fock Relativistic (HFR), and the multi-configuration Dirac-Fock (MCDF) methods. This general multiplatform approach was successfully employed in our previous studies of the K-shell spectra of Fe, O, Ne, Mg, Si, S, Ar, Ca, and Ni (e.g. Bautista et al. 2003, Garc\\'{\\i}a et al. 2005, Palmeri et al. 2003a, 2003b, 2008a, 2008b). This has the advantage that allows for consistency checks and inter-comparison. It also helps to reveal which physical processes are important for any given transition, since these different platforms employ different approaches to, for example, relativistic effects and orthogonal vs. non-orthogonal orbitals.  The present paper is organized as follows: In the next section we describe the calculations of oscillator strengths and transition rates. In Section 3 we present the calculations of collision strengths, including some new techniques specifically designed for the present problem. Our discussion of the results and conclusions are presented in Section 4.  %------------------------------------------------------------------------------- ",
        "conclusions": "We have carried out extensive calculations of transitions rates and collision strengths for electron impact excitation for the lowest 12 levels of the astrophysically important \\ion{Si}{ii} ion.  In the calculation of radiative data, we paid special attention to the weak dipole allowed and intercombination transitions, which are of particular interest for plasma diagnostics. Determination of accurate data for these transitions is particularly challenging, therefore we studied the effects of valence-valence, valence-core, and core-core interactions with three different methods, i.e. MCDF, HFR, and the central potential method implemented in {\\sc autostructure}. With MCDF we could only include valence-valence correlations, as opening of the n=2 core resulted in a large number of states that could not be managed with the computer code {\\sc grasp}. For this reason, we were able to obtain accurate transition rates for the 3s$^2$3p~$^2$P$^o$~--~3s3p$^2$~$^4$P intercombination transitions only. Both HFR and {\\sc autostructure} allowed us to investigate valence-core and core-core interaction by building very large configuration expansions. The accuracy of {\\sc autostructure} calculations were significantly improved by the use of the modified TFDA potential of \\cite{bautista08} and a new technique for optimizing the variational parameters of this potential. This optimization technique takes into account the differences between length and velocity gauges of the $gf$-values, in addition to the accuracy of predicted energy levels. Our most accurate $gf$-values were then compared with previous theoretical and experimental determinations. From these comparisons we derive a recommended set of $gf$-values and estimate their uncertainties.  We then proceed to compute electron impact excitation collision strengths with the R-matrix method. We do so by using various of the targets representations made from the previous calculations and compare the results. This allows us to identify the physical effects that affect the accuracy of the computed collision strengths. We also compare our results with those of previous calculations. The present results agree reasonably well with those of Dufton \\& Kingston (1991), who also used an orthogonal R-matrix method but with a much smaller close coupling expansion. On the other hand, the present results for Maxwellian averaged collision strengths are systematically lower than those of the recent calculation of Tayal (2008) using a non-orthogonal R-matrix method. The reasons for these differences, that typically amount to $~\\sim 30\\%$, are unclear. We argue that the source of this difference could be in the non-orthogonal R-matrix approach.   Tables 6 and 7 containing the present gf-values, A-values, and effective collision strengths are available in electronic form. The data and atomic model are also to become available through the TIPTOP\\footnote{\\tt http://heasarc.gsfc.nasa.gov/topbase} database and the XSTAR database \\cite{bautistakallman01}."
    },
    "0910/0910.5649_arXiv.txt": {
        "abstract": "{Studies of the infrared (IR) emission of cosmic sources have been proven essential to constrain the evolutionary history of cosmic star formation and that of nuclear black hole gravitational accretion, as the bulk of such events happens inside heavily dust-extinguished media. 24 $\\mu$m). } {The \\textit{Spitzer Space Telescope} has recently contributed a large set of data to constrain the nature and cosmological evolution of infrared source populations. We exploit in the present paper a large homogeneous dataset to derive a self-consistent picture of IR emission based on the time-dependent $\\lambda_{eff}$=24, 15, 12 amd 8 $\\mu$m monochromatic and bolometric IR luminosity functions (LF) over the full $0<z<2.5$ redshift range. } {Our present analysis is based on the combination of data from deep \\textit{Spitzer} surveys in the VIMOS VLT Deep Survey (VVDS-SWIRE) and GOODS areas. To our limiting flux of $S_{24}=400\\ \\mu Jy$ our derived sample in VVDS-SWIRE includes 1494 sources, and 666 and 904 sources brighter than $S_{24}=80\\ \\mu Jy$ are catalogued in GOODS-S and GOODS-N, respectively, for a total area of $\\sim$0.9 square degrees. Save for few galaxies, we obtain reliable optical identifications and redshifts, providing us a rich and robust dataset for our luminosity function determination. The final combined reliable sample includes 3029 sources, the fraction of photometric redshifts being 72\\% over all redshifts, and almost all at $z>1.5$. Based on the multi-wavelength information available in these areas, we constrain the LFs at 8, 12, 15 and 24 $\\mu$m. We also extrapolate total IR luminosities from our best-fit to the observed SEDs of each source, and use this to derive the bolometric (8-1000$\\mu$m) LF and comoving volume emissivity up to $z\\sim 2.5$. } {In the redshift interval $0<z<1$, the bolometric IR luminosity density evolves as $(1+z)^{3.8\\pm0.4}$. Although more uncertain at higher-$z$, our results show a flattening of the IR luminosity density at $z>1$. The mean redshift of the peak in the source number density shifts with luminosity: the brighest IR galaxies appear to be forming stars earlier in cosmic time ($z>1.5$), while the less luminous ones keep doing it at more recent epochs ($z\\sim1$ for $L_{IR}<10^{11}L_{\\odot}$), in general agreement with similar data in the literature. } {Our results suggest  a rapid increase of the galaxy IR comoving volume emissivity back to $z\\sim 1$ and a constant average emissivity at $z>1$. We also seem to find a difference in the evolution rate of the source number densities as a function of luminosity, a \\textit{downsizing} evolutionary pattern similar to that reported from other samples of cosmic sources. } ",
        "introduction": "\\label{intro}  A remarkable property of the infrared (IR) selected extragalactic sources is their very high rates of evolution with redshift, exceeding those measured for galaxies at other wavelengths. Indeed, based on data taken by the {\\it Infrared Astronomical Satellite} ($IRAS$),  strong evolution was detected already within the limited $0<z<0.2$ redshift range at far-IR wavelengths (60 $\\mu$m) by Hacking et al. (1987), Franceschini et al. (1988), Saunders et al. (1990). The IRAS all-sky coverage also prompted Fang et al. (1998) and Shupe et al. (1998) to established the local benchmarks at mid-IR wavelengths (12 and 25 $\\mu$m). The {\\it Infrared Space Observatory} ($ISO$) allowed us for the first time to perform sensitive surveys of distant IR sources in the mid- and far-IR, up to $z\\sim 1$ (Elbaz et al. 1999; Puget et al. 1999), showing that dust-enshrouded starbursts have undergone strong evolution both in luminosity and in density up to that redshift (e.g. Franceschini et al. 2001, Chary \\& Elbaz 2001,  Elbaz et al. 2002, Pozzi et al. 2004).  The ISO results have been further extended in redshift, to $z>1$, by the \\textit{Spitzer Space Telescope} (Papovich et al. 2004; Marleau et al. 2004; Dole et al. 2004, Lagache et al. 2004), operating between 3.6 and 160 $\\mu$m with improved sensitivity and spatial resolution compared to previous IR observatories. Le Floc'h et al. (2005), in particular, studied the evolution of IR-bright sources up to $z\\sim1$ using a sample of mid-infrared (24 $\\mu$m) sources with complete redshift information and derived the rest-frame 15 $\\mu$m and total IR LFs. Strong evolution of the IR-selected population with look-back time both in luminosity and in density was found in agreement with previous results. Similarly, Perez-Gonzalez et al. (2005) have fitted the rest-frame 12 $\\mu$m in various redshift bins in the range $0<z<3$, based on photometric redshift determinations, therefore constraining the evolution of IR-bright star-forming galaxies. Caputi et al. (2007) present the rest-frame 8 $\\mu$m LF of star-forming galaxies in two bins at $z\\sim1$ and $z\\sim2$. Babbedge et al. (2006), based on an early reduction of the SWIRE EN1 field dataset and using photometric redshifts based on relatively shallow optical imaging, computed the mid-IR (8 and 24$\\mu$m) galaxy (AGN1) LF up to $z\\sim2$ ($z\\sim4$; albeit with uncertainties strongly increasing with redshift). \\textit{Spitzer} (Werner et al. 2004) studies have also allowed accurate determinations of the $local$  {(or low-redshift}) LFs at 8 and 24 $\\mu$m (Huang et al. 2007; and Marleau et al. 2007; Vaccari et al. 2009, in preparation), respectively, essential for comparison with properties of high-$z$ sources.  All these studies indicated a strong evolution of the IR population up to $z\\sim1$ (with a comoving luminosity density produced by luminous IR galaxies more than 10 times larger at $z\\sim$1 than in the local Universe). The evolution flattens at higher redshifts, up to the highest currently probed by \\textit{Spitzer}, $z\\sim3$. Moreover, luminous and ultra-luminous IR sources, LIRGs and ULIRGs, become the dominant population contributing to the comoving infrared energy density beyond $z\\sim$0.5 and make up 70\\% of the star-forming activity at z$\\sim$1 (Le Floc'h et al. 2005).  With the aim of achieving an improved knowledge of the evolutionary properties of \\textit{Spitzer}-selected sources, we exploit in this paper the combination of data from the GOODS and VVDS-SWIRE multi-wavelength surveys to determine mid-IR and bolometric luminosity functions (and thus estimate the SFR density) over a wide redshift interval, $0<z<2.5$. The combination of data from three distinct sky areas in our analysis allows us to strongly reduce the cosmic variance effects. The coverage of the luminosity-redshift plane is also improved with the combination of surveys with two limiting fluxes, the deeper GOODS surveys (80 $\\mu$Jy limit) and the shallower wider-area VVDS-SWIRE (400 $\\mu$Jy limit) survey, with similar number of sources in the two flux regimes. The redshift information has been maximized, including a large number of new spectroscopic redshifts from public databases in the GOODS fields, and an highly optimized photometric redshift analysis in the wide-area VVDS-SWIRE field, taking advantage of the deep multi-wavelength photometry available and of the VVDS-based photometric redshift tools. The source statistics (more than 3000 objects in the combined sample) is rich enough for a tight redshift binning in the LF determination. Finally, the extremely rich suite of complementary data from the UV to the far-IR available in the three fields allows us to obtain a robust characterization of source spectra in order to compute reliable K-corrections and spectral extrapolations to be applied to our LF determination.  At the same time, the large complement of multi-band photometric data available for all our sample sources prompts us to compute reliable bolometric corrections from the rest-frame 24 $\\mu$m to the total (8-1000$\\mu$m) luminosities. These correction factors are particularly well established on consideration that the 24 $\\mu$m flux is an excellent \\textit{proxy} to the bolometric flux for active galaxies at high-redshifts (for high-$z$ sources longer wavelength fluxes are not usually available), see e.g. Caputi et al. (2007), Bavouzet et al. (2008). For these reasons we believe that our constraints on the evolutionary comoving luminosity density and star-formation rate that will be discussed in the paper are among the most solid available at the present time.    The paper is structured as follows. In Section~\\ref{SWIRE} we introduce the VVDS-SWIRE 24 $\\mu$m dataset which was first employed here. In Section~\\ref{sample} we describe the multi-wavelength identification and redshift determination process for 24 $\\mu$m sources in the GOODS and SWIRE fields. In Section \\ref{sed} we explain the procedure adopted to derive the monochromatic mid-IR rest-frame luminosities. Section 5 is devoted to the description of the computation of the LFs with the $1/V_{max}$ technique and to the presentation of our results as compared to already published data in various IR bands. The bolometric LF is also presented. Finally, in Section 6 we discuss our results about the evolution of the bolometric LF and of the IR luminosity density with redshift. Moreover, we compare our results with model predictions. A summary of the paper is presented in Section 7.  We adopt throughout a cosmology with $H_0$=70 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_M$=0.3 and $\\Omega_{\\Lambda}$=0.7. We indicate with the symbol $L_{24}$ the luminosity at 24 $\\mu$m in erg/s ($\\nu L_{\\nu}$, and similarly for other wavelengths). ",
        "conclusions": "\\subsection{The evolutionary bolometric luminosity functions} \\label{tir}  For an easier description of the evolution of the bolometric LF, we adopted a routinely used  parameterization law (Saunders et al. 1990, Pozzi et al. 2004, Le Floc'h et al. 2005, Caputi et al. 2007) to fit our data:  \\begin{eqnarray} \\lefteqn{\\Phi(L)\\,\\mathrm{d\\,log}_{10}(L)=} \\nonumber \\\\ & & \\Phi^*\\left(\\frac{L}{L^*}\\right)^{1-\\alpha}\\mathrm{exp}\\left[-\\frac{1}{2\\sigma^2}\\mathrm{log}^2_{10}\\left(1+\\frac{L}{L^*}\\right)\\right]\\mathrm{d\\,log}_{10}(L) \\end{eqnarray}  where, in this case, $L$ is the TIR [8-1000$\\mu$m] luminosity. The parameter $\\alpha$ correspond to the slope at the faint end. $L^*$ is the characteristic $L_{bol}^{IR}$ luminosity and $\\Phi^*$ is the normalization factor. The best-fitting parameters for the local LF are: $\\alpha\\simeq 1.2$, $\\sigma\\simeq 0.7$, $L^*\\simeq 1.77\\,10^9 L_{\\odot}$ and $\\Phi^*\\simeq 0.0089\\ Mpc^{-3}$.  The main uncertainties potentially affecting our results are those concerning the estimate of the bolometric luminosity from multiwavelength data limited to $\\lambda<24\\ \\mu$m, and the lack of statistics at the faint end of the LF preventing us from significantly constraining its slope. We have kept the slope of the faint-end of the LF fixed to the local observed value ($\\alpha=1.2$; see also Zheng et al. 2006, Caputi et al. 2007) and we have fitted the observations in each redshift bin by varying only $L^*$ and $\\Phi^*$, using a $\\chi^2$ minimization procedure. The slope at the bright end has been only slightly changed manually in a couple of redshift bins to provide a better fit in the highest luminosity bin. This procedure corresponds to a combination of luminosity and density evolution. The results of the fitting procedure are presented in Figure \\ref{LF_bol} (solid black lines).  From our fitting, we infer that the comoving number density of sources, as parameterized by $\\Phi^*$, evolves like ($1+z$)$^{1.1}$ in the redshift range $0<z<1$, while the characteristic infrared luminosity ($L^*$) evolves as ($1+z$)$^{2.7}$. Above $z>1$, the degeneracy between luminosity and density evolution increases, due to the more limited range in luminosity covered by our LF. However, the LFs at $z>1$ are consistent with no or limited evolution.  \\subsection{The bolometric IR luminosity density} \\label{bolirld}  One of the goals for computing the bolometric luminosity functions is to obtain an estimate of the total comoving IR luminosity density as a function of redshift. This inference has important cosmological implications, because  it is tightly related with the evolution of the comoving Star Formation Rate density and of the rate of gravitational accretion into black-holes.  Indeed, as discussed in many recent papers based mostly on \\textit{Spitzer} data (Yan et al. 2007, Sajina et al. 2007, 2008; Daddi et al. 2007; Fiore et al. 2008; see also Fadda et al. 2002, Franceschini et al. 2005), violent starburst and AGN activities appear often concomitant in \\textit{Spitzer} 24 $\\mu$m sources at $z\\sim 2$. All these analyses agree that, after excluding the most obvious type-1 AGN, it is not easy to estimate the relative AGN/starburst contributions in high-redshift sources, even including the widest SED coverage and optical/IR spectroscopic data, given the role of dust extinction in hiding the primary energy source and the very limited diagnostic power of dust re-radiated spectra.  For this reason the only attempt we have made to differentiate the AGN contribution in our estimate of the luminosity density was to account for the type-1 quasar contribution, whose power source we can entirely attribute to gravitational accretion. Although we might expect that a significant number of other among the most luminous 24 $\\mu$m sources should include some fractional AGN contributions, it is likely that the bolometric luminosity of the bulk of the population is dominated by stellar processes (e.g. Ballantyne \\& Papovich 2007). We will in any case account for the AGN-dominated sources in our later analysis.  By integrating the best-fit LF in each redshift bins, shown in Figure \\ref{LF_bol}, we obtained our determination of the total IR luminosity density up to $z\\sim 2.5$, which is reported in Figure \\ref{rho} as blue filled squares and tabulated in Table \\ref{dens_values}. In addition to the Poisson noise, the errorbars that we associated to our estimates account for two major sources of uncertainties: the slope of the faint end of the luminosity function (which was set to $\\alpha=1.2$ with an error inferred from fitting the local LF), and the error on the normalization parameter ($\\Phi^*$) in the $\\chi^2$ minimization procedure.  Within the previously mentioned uncertainties, we find a general agreement with previously published results based on IR data. The orange filled region in Figure \\ref{rho} shows the results of Le Floc'h et al. (2005) up to $z\\sim1$, the light pink area marks the data obtained by Caputi et al. (2007) and the magenta curve those by Perez-Gonzalez et al. (2005). In the redshift range $0<z<1$, we find that the total bolometric IR luminosity density evolves as $(1+z)^{3.8\\pm0.4}$. This evolution is very close to that derived by Le Floc'h et al. (2005), who found $(1+z)^{3.9}$ between $z=0$ and $z=1$, but is significantly steeper than estimated by Caputi et al. (2007), $(1+z)^{3.1}$. At higher redshifts our results suggest  a flattening of the IR luminosity density above $z\\sim 1$. The results from the IR data are also in fully agreement within what is derived from the far-ultraviolet data: a strong increase up to $z\\sim1$, and a flatting above as shown for example in Tresse et al. (2007).  We also report in Fig. \\ref{rho} the separate contributions from LIRGs (here defined as objects with $10^{11}<L_{IR}<10^{12}L_{\\odot}$, green dashed line) and ULIRGs ($L_{IR}>10^{12}L_{\\odot}$, red dot-dashed line) to the IR luminosity density. At $z\\sim0.2$, $\\sim$30\\%  of the bolometric IR luminosity density is contained in LIRGs and $<1$\\% in ULIRGs. At $z=1$, we find that LIRGs and ULIRGs contribute 45\\% and 6\\%, respectively, to the total IR luminosity density. At $z\\sim2$, the contributions of LIRGs and ULIRGs become 48\\% and 45\\% of the total budget, respectively: the cosmic evolution of the highest- and of the moderate-luminosity sources appear to be drastically different. This finding is consistent with the results of previous \\textit{Spitzer} studies of the specific SF for star-forming galaxies (e.g. Perez-Gonzalez et al. 2005, P. Santini et al. 2009).  Finally, we have estimated the comoving luminosity densities after having taken out the contribution of the 180 sources identified as type-1 AGNs from our SED fitting, as explained in Sect. \\ref{sed}. However, this does not make any significant difference in Fig. \\ref{rho}. Assuming that the remaining part of the population has a bolometric emission dominated by star-formation, we report on the right-hand ordinate axis of Fig. \\ref{rho} the comoving SFR density, that we calculate from Kennicutt et al. (1998): $SFR [M_{\\odot}/yr] = 1.7 \\times 10^{-10} L([8-1000\\mu m])/L_{\\odot}$ .  \\subsection{Evolution of the comoving source space density}  A more detailed inspection of the evolutionary properties of the IR selected sources is presented in Figure \\ref{dens} in terms of the variation with redshift of the comoving source space densities for different IR luminosity classes. The data points here have been derived from the best-fit to the bolometric LFs shown in Figure \\ref{LF_bol}, but we report only the data within luminosity intervals acceptably constrained by our observations.  The figure indicates a systematic shift with redshift of the number density peak as a function of luminosity: the brighest IR galaxies formed earlier in the cosmic history ($z> \\sim1.5$), while the number denisty of the less luminous ones peaks at lower redshifts ($z\\sim1$ for $L_{IR}\\leq 10^{11}$). Although our results do not constrain the high-redshift evolution of the lower luminosity sources, they provide a tentative detection of a $downsizing$ effect in the evolution of the IR galaxy populations, hence in the cosmic star formation history. This behaviour appears to parallel a similar one reported for a completely different population of cosmic sources, i.e. type-1 AGNs selected in X-rays and analyzed by Hasinger et al. (2005, see also Smolcic et al. 2009 for radio selected galaxies).  If confirmed, our results are obviously much more significant and further reaching, because they concern the bulk of the emission by the whole cosmic source population selected in the IR. In any case, this similarity of the evolutionary properties of our IR-selected galaxy population and those of X-ray AGNs strongly supports the case for co-eval AGN/starburst activity and for a physical relation between the two populations.   \\begin{figure}[!h] \\includegraphics[width=9.5cm,height=10cm]{dens_downs1_noAGN.ps} \\caption{The space density of IR selected galaxies as a function of redshift in different luminosity classes. The numbers with the same colour coding at the left side of each curve indicate the corresponding IR luminosity in solar units (logarithmic). } \\label{dens} \\end{figure}  \\subsection{Comparison with model predictions} \\label{AF}  We report in Figures \\ref{zdist-080-muJy}, \\ref{mips24-dif-counts}, \\ref{8mu}-\\ref{24mu} and \\ref{rho} a comparison of our results with predictions of a phenomenological model by Franceschini et al. (2009,  AF09). The latter, in particular, are compared with the rest-frame 24 $\\mu$m LF in Figure \\ref{24mu} (dashed lines), showing a generally good agreement. The low-$z$ behaviour of the AF09 model has been calibrated on a variety of number count data at bright fluxes (including the IRAS all-sky IR counts and those from the \\textit{Spitzer} SWIRE project), which guarantees excellent control of the low-redshift universal emissivity in the IR.  Compared to previous analyses essentially based on pre-\\textit{Spitzer} surveys (e.g. Franceschini et al. 2001), an important addition in AF09 was the introduction of an evolutionary component of very luminous IR galaxies dominating the cosmic IR emissivity at $z\\sim$2, and essentially absent or very rare locally, hence caracterized by an extremely fast evolution in cosmic time. This population is revealed from the analysis of the galaxy samples selected by the \\textit{Spitzer}/MIPS deep 24$\\mu$m and SCUBA sub-millimetric observations, while it was undetected by ISO.  In Figure \\ref{rho} we compare the redshift evolution of our derived cosmic star-formation rate with the AF09 model. The general trend is nicely reproduced, implying (together with all the other observational constraints considered by the model) that the main features of the IR galaxy evolution are understood. Our results support a scenario where the fast evolution observed from $z=0$ to 1 flattens above $z\\sim 1$ and keeps approximately flat up to $z\\sim2.5$. Moreover, both models and the present data provide clear evidence for the existence of a population of very luminous galaxies becoming increasingly important at $z>1$. The current interpretation identifies these objects with the progenitors of the spheroidal galaxies. A more detailed discussion is deferred to AF09."
    },
    "0910/0910.0546_arXiv.txt": {
        "abstract": "We study the compact binary population in star clusters, focusing on binaries containing black holes, using a self-consistent Monte Carlo treatment of dynamics and full stellar evolution.  We find that the black holes experience strong mass segregation and become centrally concentrated.  In the core the black holes interact strongly with each other and black hole-black hole binaries are formed very efficiently.  The strong interactions,  however, also destroy or eject the black hole-black hole binaries.  We find no black hole-black hole mergers within our simulations but produce many hard escapers that will merge in the galactic field within a Hubble time.  We also find several highly eccentric black hole-black hole binaries that are potential LISA sources, suggesting that star clusters are interesting targets for space-based detectors.  We conclude that star clusters must be taken into account when predicting compact binary population statistics. ",
        "introduction": "\\label{intro}  The inspirals and mergers of compact binaries where both members are neutron stars (NS) or black holes (BH) are some of the most promising sources for the current and next generation of ground based gravitational wave (GW) detectors (LIGO, Virgo, GEO 600, TAMA 300).  NS-NS binaries are expected to be the most plentiful merger species in the frequency regime of ground based detectors \\citep{AbbottNS,Belczynski07}, however BH-BH binaries are more massive and thus can be detected at larger distances, as far as the Virgo cluster \\citep{AbbottBH} for the current generation of gravitational wave detectors and up to cosmological distances for the next generation.  Compact binaries with longer periods and including white dwarfs (WD) may also be detectable in the low-frequency ($5 \\times 10^{-5} - 1$ Hz) LISA (Laser Interferometer Space Antenna) frequency band \\citep{HBW90,Benacquista01,Nelemans01,BBandB08}.  WD-WD binaries will be the most plentiful stellar mass LISA sources and are expected to produce confusion limited noise (e.g. \\citealt{EIS87,HBW90,Nelemans01,Timpano06,Ruiter08}) whereas the less common NS-NS, NS-BH, and BH-BH binaries are potentially resolvable.  In order to make predictions for GW event rates, the population of compact binaries in the universe must be understood.  While the population of NSs in the local universe can be constrained by observations of pulsars (e.g. \\citealt{Kalogera01,Lorimer05}), BHs cannot be observed directly and their properties can only be constrained by modelling.  There have been several studies carried out on the compact binary population in the galactic field where stellar and binary evolution proceeds in isolation and can be modelled using simple population synthesis.  In particular \\cite{Belczynski07} predict a detection rate for the advanced LIGO detector of $\\sim 20$ NS-NS mergers yr$^{-1}$, only $\\sim 2$ BH-BH mergers yr$^{-1}$, and $\\sim 1$ NS-BH merger yr$^{-1}$.  Thus the detection rate in the galactic field should be dominated by NS-NS mergers.  \\cite{BBandB08} have performed similar calculations for the galactic field in the LISA band.  Depending upon the assumptions made about the probability of mergers during common envelope evolution they find $2-6$ resolvable NS-NS binaries and $0-5$ resolvable BH-BH binaries.  This implies that, although rare, both NS-NS and BH-BH binaries may appear as resolved stellar mass sources in the LISA band.  Overall NS-NS binaries will dominate the field detection rate for ground-based detectors while a detection in the LISA band is possible but unlikely.  In star clusters the situation is rather different.  Here interactions between stars and binaries are common (e.g. \\citealt{Heggie75}) and can affect the final outcome of binary evolution.  Such interactions can form new binaries from single stars, exchange binary members and field stars and can reduce or increase the period of existing binaries.  In order to maintain energy equipartition, interactions between particles of different masses tend to accelerate the lowest-mass particles to the highest velocities \\citep{Spitzer87}.  As a consequence low-mass objects are the most likely to escape during few-body encounters.  Therefore few-body interactions tend to introduce massive objects into binaries.  Massive objects also tend to sink to the center of star clusters where the stellar density is highest (mass segregation, also a consequence of energy equipartion, \\citealt{Spitzer87}) and thus massive objects are the most likely to experience dynamical interactions.  Since BHs rapidly become the most massive objects in star clusters due to stellar evolution they will be particularly strongly affected by dense stellar environments and are very likely to be exchanged into binaries.  Therefore star clusters are predicted to form BH-BH binaries rather efficiently \\citep{SigPhin93} and may significantly enhance the BH-BH merger rate in the universe.  Several authors have investigated this possibility using various approximations.  \\cite{Gultekin04} and \\cite{Oleary06} have simulated the formation of intermediate-mass black holes (IMBHs) assuming that the BHs in the cluster are completely mass-segregated and interact only with each other.  In this situation, where the BHs interact very strongly, the authors resolve these encounters with explicit few-body integration.  The interactions can lead to the efficient formation of BH-BH binaries but also tend to destroy or eject those already formed.  These authors conclude that gravitational wave mergers can occur within young star clusters but BH-BH binaries will be destroyed rather efficiently and the population will be depleted within a few Gyrs.  \\cite{Oleary06} in particular calculate a star cluster BH-BH detection rate for advanced LIGO of $1-10$ merges yr$^{-1}$, up to 70\\% of which actually occur in ejected BH-BH binaries and thus take place in the galactic field.  Neither set of authors include stellar evolution in their simulations.  By contrast \\cite{Iva08} and \\cite{Sadowski08} have conducted studies of the compact binary population in star clusters using stellar evolution prescriptions and simplified two-zone models of cluster dynamics.  The two-zone models assume that the BHs and BH-BH binaries remain in dynamical equilibrium with the rest of the cluster and do not strongly mass-segregate.  Thus the density of BHs and BH-BH binaries remains low and they do not interact with each other nearly as frequently as in the previous models.  \\cite{Sadowski08} in particular find no NS-NS mergers but a much higher rate of BH-BH mergers than \\cite{Oleary06}.  They calculate a detection rate of $25-3000$ mergers yr$^{-1}$ for advanced LIGO even though their treatment of the few-body interactions is the same as for \\cite{Oleary06}.  This is because although there are fewer interactions that create BH-BH binaries there are also fewer interactions that destroy them.  This highlights the importance assumptions about global cluster dynamics can have on detection rates.  The only study with full treatment of both dynamics and stellar evolution is that of \\cite{SPZMcMill00} who conducted small ($N \\sim$ a few $10^{3}$) direct N-body simulations and showed that BH-BH binaries are quickly ejected from star clusters due to strong few-body interactions.  This seems to confirm the model of \\cite{Oleary06} but the simulations are too small for any strong conclusions to be drawn.  It seems that star clusters can significantly enhance the rate of BH-BH mergers in the universe however by exactly how much depends strongly on the dynamical assumptions made.  All of these simulations resolve the few-body interactions but either use very simplified models for the global cluster dynamics or have values of $N$ too small to represent GCs.  In this paper we use a Monte Carlo code to self consistently model the dynamics of globular clusters over a range of metallicities, binary fractions, and initial concentrations.  We hope to constrain which dynamical assumptions are most likely to produce accurate results for gravitational wave event rate predictions.  We will focus only on the compact binaries that can be found within the cluster during its evolution, leaving a detailed discussion of the population of escapers to a future paper.  In \\S~\\ref{methods} we briefly describe the Monte Carlo code and some of its features and limitations.  In \\S~\\ref{IC} we describe our initial models.  In \\S~\\ref{results} we present the results of our simulations.  In \\S~\\ref{GWdetect} we present predictions for LISA detections.  We discuss our results in \\S~\\ref{discussion} and conclude in \\S~\\ref{conclude}. ",
        "conclusions": "\\label{conclude}  We have studied the dynamics of the BH population in star clusters with a self-consistent Monte Carlo treatment of the global dynamics and full stellar and binary evolution.  We confirm the predictions of \\cite{SigPhin93} that BH-BH, but not NS-NS or NS-BH, binaries are produced efficiently in star clusters by dynamical interactions.  We find that the BHs are mass-segregated and interact strongly with each other, confirming the more approximate models of \\cite{Oleary06}.  We find no BH-BH mergers within the clusters we simulate but many hard BH-BH escapers that will merge in the galactic field within a Hubble time.  Detection rates for both ground- and space-based detectors will be calculated and presented in a future paper.  We find that our simulations produce potential LISA sources that, while rare, will be highly eccentric and may still represent a significant enhancement to the galactic field population.  We will certainly produce more detections when the escapers are included.  We conclude that star clusters produce BH-BH binaries efficiently and must be taken into account when considering detection rates for both ground- and space-based detectors."
    },
    "0910/0910.2665_arXiv.txt": {
        "abstract": "Over the last number of years spectroscopic studies have strongly supported the assertion that protostellar accretion and outflow activity persists to the lowest masses. Indeed, previous to this work the existence of three brown dwarf (BD) outflows had been confirmed by us. In this paper we present the results of our latest investigation of BD outflow activity and report on the discovery of two new outflows. Observations to date have concentrated on studying the forbidden emission line (FEL) regions of young BDs and in all cases data has been collected using the UV-Visual Echelle Spectrometer (UVES) on the ESO VLT. Offsets in the FEL regions are recovered using spectro-astrometry. Here ISO-Oph 32 is shown to drive a blue-shifted outflow with a radial velocity of 10-20 \\km\\ and spectro-astrometric analysis constrains the position angle of this outflow to 240$^{\\circ}$ $\\pm$ 7$^{\\circ}$. The BD candidate \\ISO is found to have a bipolar outflow bright in several key forbidden lines (V$_{RAD}$ = -20 \\km, +40\\km) and with a PA of 190$^{\\circ}$-210$^{\\circ}$. A striking feature of the \\ISO outflow is the strong asymmetry between the red and blue-shifted lobes. This asymmetry is revealed in the relative brightness of the two lobes (red-shifted lobe is brighter), the factor of two difference in radial velocity (the red-shifted lobe is faster) and the difference in the electron density (again higher in the red lobe). Such asymmetries are common in jets from low mass protostars (0.1 \\Msun to 2 \\Msun) and the observation of  a marked asymmetry at such a low mass ( $<$ 0.1 \\Msun) supports the idea that BD outflow activity is scaled down from low mass protostellar activity. Also note that although asymmetries are unexceptional, it is uncommon for the red-shifted lobe to be the brightest as some obscuration by the accretion disk is assumed. This phenomenon has only been observed in one other source, the classical T Tauri (CTTS) star RW Aur. The physical mechanism responsible for the brightening of the red-shifted lobe although as yet unknown must also now apply at BD masses to include the \\ISO\\ outflow. In addition to presenting these new results, a comprehensive comparison is made between BD outflow activity and jets launched by CTTSs.  In particular, the application of current methods for investigating the excitation conditions and mass loss rates ($\\dot{M}_{out}$) in CTT jets to BD spectra is explored. Where possible the Bacciotti $\\&$ Eisl{\\\"o}ffel technique is used to study the ionisation fraction, electron temperature and total density. For LS-RCrA 1, \\ISO and ISO-Oph 102 $\\dot{M}_{out}$ is measured to be in the range 10$^{-10}$ to 10$^{-9}$ \\Msun yr$^{-1}$ using a method based on the luminosity of the [OI]$\\lambda$6300 and [SII]$\\lambda$6731 lines. Mass loss rates for our sample of BD outflows are found to be comparable to the mass accretion rates. Overall, as our results and discussion show, what is currently known about outflow activity in the BD mass regime points to strong similarities between BD outflows and CTT jets. ",
        "introduction": "The outflow phenomenon, much studied in young stars has been observed in a wide variety of astrophysical objects from low mass protostars to active galactic nuclei (AGN) \\citep{Reipurth01, Camenzind05}. Brown Dwarfs (BDs) are the newest objects to be added to this list \\citep{Fernandez01, Barrado04, Whelan05, Whelan07, Phan08, Whelan09}. We are currently leading a project to search for and study outflows driven by young BDs and very low mass stars. The overall aim of this work is to explore outflow activity in the BD mass regime through comparison with observations of outflows driven by low mass protostars and the classical T Tauri stars (CTTSs) in particular. Much emphasis has been placed on understanding how accretion activity \\citep{Gatti06, Herczeg09} and planet formation processes \\citep{Pasucci09} differ between low mass protostars and young BDs. However, little is still known about whether the outflow mechanism varies as mass decreases towards the BD regime. Observations aimed at detecting outflows from such low mass objects are difficult and rare and to date six candidates have been investigated by us spectro-scopically at visible wavelengths.    Outflows from low mass protostars are revealed in several forms, from slow wide angled winds to molecular outflows and jets \\citep{Reipurth01}. A protostellar jet is a relatively fast (typical radial velocity = 200 \\km) collimated outflow normally seen within a few 1000 AU of a young star and bright in optical and/or near-infrared shock tracers such as [SII]$\\lambda$6731, [FeII] at 1.644$\\mu$m or H$_{2}$ at 2.122 $\\mu$m. CTTSs are Class II low mass protostars \\citep{Lada87, Andre93}, meaning that they are coming towards the end of their accretion/outflow phase and have crossed the so-called ``birth-line\" \\citep{Stahler83}. As they are optically visible the jet can be traced very close to the protostar, hence, they offer a unique opportunity to study the launching of protostellar jets at high spatial resolution \\citep{Ray07}. Jets from CTTSs are traditionally probed at forbidden emission line (FEL) wavelengths and long-slit and integral field spectroscopic techniques have been hugely important in their study \\citep{Dougados00, Whelan04}. That CTT-like outflows could be launched by actively accreting BDs was first suggested when high quality spectra revealed the presence of FEL regions in the optical spectra of BDs known to be accretors \\citep{Fernandez01}. While the FEL regions were easily identified they were considerably fainter than those detected in the spectra of CTTSs \\citep{Masciadri04} and any extension in the form of an outflow was not apparent \\citep{Fernandez01, Fernandez05}. Hence their origin in an outflow could not be established. The critical density region for FELs occurs much closer to the driving source of the outflow for nearby ($\\sim$ 150 pc) young BDs than for low mass protostars and for BDs is estimated to be on scales of $\\sim$ 100 mas \\citep{Whelan05}. As the most intense forbidden emission coincides with the critical density region it is currently challenging to directly resolve a BD outflow in FELs and this goal has not yet been achieved. Our approach to this problem has been to obtain high quality spectra using the UV-Visual Echelle Spectrometer (UVES) on the European Southern Observatory's (ESO) Very Large Telescope (VLT) and to recover the spatial offset in the region of forbidden emission using spectro-astrometry.       The sources targeted so far meet at least one of two criteria. Firstly they all have masses at or significantly below the hydrogen burning mass limit (HBML) and secondly they are known to be powerful accretors. Their classification as a strong accretor is based on studies of their \\Ha\\ lines \\citep{Jay03a} and on measurements of their mass accretion rates (see Table 6). In the vast majority of cases the mass of young BD candidates cannot be directly measured and estimates are based on theoretical predictions \\citep{Luhman00} and/or spectral modelling \\citep{Mohanty04}. % The objects investigated by us so far are listed in Table \\ref{tab1} along with their theoretical masses and spectral types. Previous to this paper we have reported on the discovery of and discussed the optical outflows from, ISO-Oph 102 \\citep{Whelan05}, 2MASS1207-3932 \\citep{Whelan07} and \\ls \\citep{Whelan09}. With a mass of only 24 M$_{JUP}$ \\MASS\\ is at present the lowest mass galactic object driving an outflow. In this paper we present the results of our recent observations of ISO-ChaI 217, DENIS-P J160603.9-205644 and ISO-Oph 32 and report for the first time outflows driven by \\ISO and ISO-Oph 32. Of the six objects observed by us to date only one, \\ISO\\ has a theoretical mass which is just above the HBML, placing it at present just outside the mass range for a BD. However,  with a predicted mass of just 80 \\J\\ it is still one of the lowest mass objects known to launch an outflow. Therefore for the purpose of the discussion we will include it in this paper with the likely BDs. In addition to presenting our spectro-astrometric analysis, a technique for constraining the position angle (PA) of the BD outflows is outlined, methods for studying the physical conditions of the BD outflow and of estimating the mass outflow rate are discussed and similarities between BD outflows and CTT jets highlighted and considered. ",
        "conclusions": "Currently our understanding of astrophysical outflows is focussed on observations of low mass protostellar objects. Although much can be gleaned from such observations a proper appreciation of the outflow phenomenon can only be reached through a complete investigation of how outflow activity compares for a range of driving sources, varying in mass and age. The new results included in this paper include the presentation of two new BD outflows, the unusual red-shifted asymmetry in the \\ISO\\ and the demonstration that  $\\dot{M}_{out}$ in BDs is comparable to $\\dot{M}_{acc}$. Overall, the main aim of our BD project is to explore how substellar outflow activity compares to what is well documented for low mass protostars. As our sources are all optically visible our comparison is focussed on the CTTSs. A second advantage of this work is that it may add to the debate on the formation mechanism of BDs i.e. whether they form like low mass protostars or are ejected embryos from multiple systems. The later mechanism is thought to produce truncated disks and very little envelope. If BD outflows are long-lived this would argue against such a scenario although further investigation is required.  Due to the faintness of the BD outflows observations are difficult and time consuming. However data gathered to date has led to some significant results. The most compelling pieces of evidence which support the continuation of CTT-like outflow activity into the BD mass regime include, the detection of an outflow component to the \\ls \\Ha\\ emission line and the suggestion of a possible disk dust hole, the identification of an LVC to the FELs of \\ls and ISO-Oph 32, the strong asymmetry in the \\ISO outflow and the discovery of a molecular component to the optical outflows of \\ls and ISO-Oph 102.  While these results are significant much further work is needed to provide a more complete picture of how BD and CTT outflows compare. This is especially true for any study of the physical conditions in BD outflows and BD mass outflow rates. A considerably increased sample is needed. This can be provided through further optical and NIR spectroscopic studies complemented by millimetre studies.  Investigations must be extended to other wavelength regimes so that NIR atomic and molecular outflow tracers can be identified and further CO outflows found. Finally optical/NIR images would allow the collimation of the BD outflows to be probed. Many of these projects are under-way or planned however due to faintness of the emission from a typical BD outflows, it is likely that long-slit/echelle spectroscopy combined with spectro-astrometry will remain (in the short-term at least) the most effective way to investigate BD outflow activity.  \\begin{center} \\begin{figure} \\includegraphics[width=9.5cm]{ISO_217OI30_0.pdf} \\includegraphics[width=9.5cm]{ISO_217OI30_90.pdf} \\caption{Spectro-astrometric analysis of the ISO-Cha217 \\Ox\\ line. The line has a blue and a red-shifted component. In the spectrum taken at 90$^{\\circ}$ the blue-shifted component to the line appears as an extended wing. Offsets decrease at 90$^{\\circ}$ and comparison with measurements at 0$^{\\circ}$ constrain the outflow PA at $\\sim$ 208$^{\\circ}$. Offsets are a moving sum across the line and continuum position. The displacement in the line is measured in the continuum subtracted spectrum. The 1-$\\sigma$ error in the line is marked by the dashed line. Finally as outlined in section 2 the line and continuum emission is summed in such a way that the 1-$\\sigma$ error in both is comparable.} \\end{figure} \\end{center}       \\begin{center} \\begin{figure} \\includegraphics[width=9.5cm]{ISO_217OI63_0.pdf} \\includegraphics[width=9.5cm]{ISO_217OI63_90.pdf} \\caption{Spectro-astrometric analysis of the ISO-Cha217 [OI]$\\lambda$6363 line. Analogous to the \\Ox\\ line, the line is made of blue and red-shifted emission which spectro-astrometry shows to be offset in opposing directions. As expected the offsets are similar to what is measured for the \\Ox\\ line. Again comparison between measurements at 0$^{\\circ}$ and 90$^{\\circ}$ constrain the outflow PA at $\\sim$ 205$^{\\circ}$. As for all the ISO-ChaI 217 FELs, offsets are a moving sum across the continuum and line position and the 1-$\\sigma$ error in both measurements is comparable. } \\end{figure} \\end{center}  \\begin{center} \\begin{figure} \\includegraphics[width=9.5cm]{ISO_217Ha_0.pdf} \\includegraphics[width=9.5cm]{ISO_217Ha_90.pdf} \\caption{Spectro-astrometric analysis of the ISO-Cha1 217 \\Ha\\ line. Analysis was carried out in the same manner, as outlined for the \\OI\\ lines. No signal from the outflow is detected in the \\Ha\\ line. Emission is coincident with the source hence it is reasonable to assume that the bulk of the emission is coming from the accretion flow. For the \\Ha\\ line the spectro-astrometric accuracy achieved at the line peak is much greater than in the wings or for the the continuum emission. Hence for all the \\Ha\\ lines emission is summed so that the error in the line-wing and continuum is comparable. } \\end{figure} \\end{center}   \\begin{center} \\begin{figure} \\includegraphics[width=9.5cm]{ISO_217SII16_0.pdf} \\includegraphics[width=9.5cm]{ISO_217SII16_90.pdf} \\caption{Spectro-astrometric analysis of the ISO-Cha1 217 [SII]$\\lambda$6716 line. Again the line was observed at a PA of 0$^{\\circ}$ and 90$^{\\circ}$ degrees. The blue and red-shifted emission is more clearly separated spectrally than in the \\OI\\ lines. In addition it is offset further from the BD position. Analysis of the [SII]$\\lambda$6716 line constrains the outflow PA at $\\sim$  193$^{\\circ}$. This is smaller than but similar to the estimate from the \\OI\\ lines. Note the asymmetry in the red and blue lobes of the outflow in both velocity and intensity. } \\end{figure} \\end{center}   \\begin{center} \\begin{figure} \\includegraphics[width=9.5cm]{ISO_217SII_0.pdf} \\includegraphics[width=9.5cm]{ISO_217SII_90.pdf} \\caption{Spectro-astrometric analysis of the ISO-Cha1 217 [SII]$\\lambda$6731 line. Analysis agrees with results for the [SII]$\\lambda$6716 line and constrains the outflow PA at $\\sim$ 200$^{\\circ}$. Again the red and blue lobes of the outflow are clearly separated and the asymmetry is especially clear. The difference between the intensity of the red and blue lobes as traced by the [SII]$\\lambda$6731 line at 0$^{\\circ}$ is striking.} \\end{figure} \\end{center}   \\begin{center} \\begin{figure} \\includegraphics[width=11.5cm]{SII_PV.pdf} \\caption{Position velocity diagrams of the \\ISO\\ [SII]$\\lambda$6731 line region, comparing the line at slit PAs of 0$^{\\circ}$ (left) and at 90$^{\\circ}$ (right). For the PV diagrams the emission line region was smoothed using an elliptical Gaussian filter as described above. Contours begin at 3 times the r.m.s noise and increase by factors of $\\sqrt{2}$. The black line delinates the BD position. The PV plots clearly support the spectro-astrometric results. The relative offset between the red and blue-shifted part of the line can be seen by eye in the plot for 0$^{\\circ}$. As this offset is a factor of 4 smaller at 90$^{\\circ}$ it is not obvious in the corresponding PV diagram. A difference in brightness between the red-shifted components is also apparent with the red-shifted component at 90$^{\\circ}$ appearing marginally brighter. As we are only detecting emission very close to ISO-Cha1 217 and given that our slit width is 1 \\arcsec we would not expect to see any significant variation in brightness with PA. The apparent difference here is due to a change in observing conditions between the acquisition of the two spectra with the seeing decreasing from 0\\farcs85 for the 0$^{\\circ}$ spectrum to 0\\farcs7 at 90$^{\\circ}$. We measure the integrated flux of the red-shifted components to be comparable with the value being 0.2, 0.45 $\\times$ 10$^{-16}$ ergs/s/cm$^{2}$/A and 0.21, 0.5 $\\times$ 10$^{-16}$ ergs/s/cm$^2$/A for the red/blue-shifted components at 0/90 degrees respectively.} \\end{figure} \\end{center}    \\begin{center} \\begin{figure} \\includegraphics[width=9.5cm]{Oph32_OI_0.pdf} \\includegraphics[width=9.5cm]{Oph32_OI_90.pdf} \\caption{\\Ox\\ line emission and related spectro-astrometric analysis for the young BD ISO-Oph 32. As reported for other BDs, the \\Ox\\ line was the only FEL found in the spectrum strong enough for spectro-astrometric analysis. Spectro-astrometry confirms the origin of the line in an outflow. The measured offset in the line is greater at 90$^{\\circ}$ than at 0$^{\\circ}$ and from this comparison the outflow PA is estimated at 240$^{\\circ}$. In the case of ISO-Oph 32 the line and continuum regions were smoothed rather than summed.} \\end{figure} \\end{center}  \\begin{center} \\begin{figure} \\includegraphics[width=9.5cm]{Oph32_Ha_0.pdf} \\includegraphics[width=9.5cm]{Oph32_Ha_90.pdf} \\caption{Spectro-astrometry of the ISO-Oph 32 \\Ha\\ line. As in the case of \\ISO\\ offsets are a moving sum across the line and continuum position and the line and continuum are summed in such as way that the error in the line-wings and continuum is comparable. The $\\pm$ 1-$\\sigma$ error is delinated by the dashed lines. As for all the sources discussed in this paper no signature of outflow activity is detected in the \\Ha\\ line.} \\end{figure} \\end{center}  \\begin{center} \\begin{figure} \\includegraphics[width=9.5cm]{DENIS_Ha_0.pdf} \\includegraphics[width=9.5cm]{DENIS_Ha_90.pdf} \\caption{The \\Ha\\ line and spectro-astrometric analysis for DENIS-P J160603.9-205644. DENIS-P J160603.9-205644 is known to be actively accreting, however no evidence of outflow activity is found in our spectra. As is clear from the analysis shown here the \\Ha\\ emission is coincident with the source position and therefore must originate very close to the source. This analysis was carried out in the same manner as for all the other \\Ha\\ lines. } \\end{figure} \\end{center}"
    },
    "0910/0910.2723_arXiv.txt": {
        "abstract": "% Two possible explanations for the type SNe Ia supernovae observations are a nonlinear, underdense void embedded in a matter-dominated Einstein-de Sitter spacetime or dark energy in the $\\Lambda$CDM model. Both of these alternatives are faced with Copernican fine-tuning problems. A case is made for the void scenario that avoids introducing undetected dark energy. ",
        "introduction": "The standard $\\Lambda$CDM model is based on Einstein's gravitational theory (GR) and the cosmological principle: there is no special place in the universe - the universe is isotropic and homogeneous. The assumptions of an isotropic and homogeneous universe were used to develop the Lema\\^{i}tre-Friedmann-Robertson-Walker (LFRW) expanding universe model, which forms the basis of the standard modern $\\Lambda$CDM cosmology. Although this model has successfully described the available cosmological observations, in particular, the WMAP CMB data~\\cite{Komatsu}, it has about 8 free adjustable parameters and relies heavily on the simplest assumption of maximal symmetry, the LFRW geometry of spacetime and that GR is the correct theory of gravity to describe the large scale structure of the universe.  A disturbing feature of the standard $\\Lambda$CDM model is that roughly 96 percent of the universe is invisible. Approximately 30 percent is composed of ``dark matter\", while 70 percent is the dark energy that pervades the universe like a modern-day ether, and is responsible for the asserted acceleration of the expansion of the universe invoked to explain the supernovae observations~\\cite{Perlmutter,Riess,Kowalski}. The dark matter has been postulated to explain the rotation curve data of galaxies, the stability of clusters of galaxies, merging clusters such as the Bullet Cluster, gravitational lensing and the WMAP data. After several years of searching for the ubiquitous dark matter, no successful detection of dark matter particles has been achieved. By its nature the dark matter particles interact only with gravity, so the existence of dark matter is inferred through gravitational observations such as galaxy rotation curves.  The dark energy is a mysterious and undetectable uniform substance that has been identified with negative pressure vacuum energy. This interpretation falls prey to the serious lack of understanding of vacuum density in quantum field theory, leading to preposterous degrees of fine-tuning to agree with the ``observed\" vacuum density associated with the cosmological constant $\\Lambda$. To avoid this problem, many modified gravity theories have been proposed that devote themselves to explaining the accelerated expansion of the universe. Many of these theories suffer from maladies that render them unphysical. They can possess ghosts and instabilities and not be able to explain the precise relativistic corrections engendered by GR in the solar system, such as the Cassini spacecraft observation of the Eddington-Robertson PPN parameter $\\gamma -1= (2.1\\pm 2.3)\\times 10^{-5}$~\\cite{Bertotti}.  An early extensive study of cosmological voids~\\cite{Moffat,Moffat2} in the exact Lema\\^{i}tre-Tolman-Bondi (LTB) inhomogeneous solution of GR~\\cite{Lemaitre,Tolman,Bondi,Bonnor,Moffat3,Moffat4,Krasinski}, predicted that the luminosity distance in a void solution would depend on the red shift $z$ in a way that deviates from the LFRW model prediction~\\cite{Moffat,Moffat2}. The distance modulus calculated from the void model in 1994-95 suggested an {\\it apparent} acceleration of the expansion of the universe as inferred by an observer in an FLRW spacetime~\\cite{Moffat2}. Indeed, four years before the celebrated supernovae SNe Ia observations, it was clear that {\\it if the inhomogeneous void scenario is correct}, then the observed dimming of the supernovae light was to be expected. However, the void solution demanded that the observer be situated close to the center of the void to maintain the observed isotropy of the large scale structure of the universe. This would imply a violation of the cosmological Copernican Principle. The adherence to a dark energy such as the vacuum energy in the $\\Lambda$CDM model, also suffers from an anti-Copernican principle fine-tuning in coordinate time. The negative pressure dark energy and the associated acceleration of the expansion of the universe began dominating the evolution of the standard model when life first appeared on our planet. This is referred to as the ``coincidence\" problem. Thus, we are faced with having to make a choice between two theoretical evils, the dark energy scenario or the more conservative inhomogeneous void scenario. As has been recently demonstrated~\\cite{Bassett,Clifton,Caldwell,Clarkson,Tomita}, it is difficult with the presently available observational data to distinguish between the idealistic Copernican Principle cosmology and those cosmological models that violate it.  One ongoing vexing problem with the postulate of a cosmological constant in the $\\Lambda$CDM standard model is our lack of understanding of the vacuum density in physics. The notorious cosmological constant problem remains unresolved, leading to a huge fine-tuning problem. For example, in the standard electroweak model with a Higgs particle the predicted vacuum density is of order $10^{56}$ times bigger than the vacuum density required in the $\\Lambda$CDM model~\\cite{Higgs}. The problem becomes even more pronounced when energy scales reach the Planck energy $\\sim 10^{19}$ GeV, resulting in a fine-tuning of order $10^{122}$ compared to the ``observed\" cosmological constant.  A relativistic modified gravity (MOG) theory known as the Scalar-Tensor-Vector Gravity (STVG) theory has been applied successfully to fit astrophysical as well as cosmological data~\\cite{Moffat5,Moffat6,MoffToth,MoffToth2,MoffToth3,MoffToth4,MoffToth5,Brownstein} without non-baryonic dark matter. Thus, if exotic dark matter remains undetected (as was the nineteenth century ether), then to explain the extensive astrophysical and cosmological data, we are forced to modify Newtonian and Einstein gravity. In~\\cite{MoffToth2,MoffToth3}, the MOG was able to also explain the accelerated expansion of the universe without a cosmological constant, although this explanation was also subject to the fine-tuned coincidence problem.  In contrast to this need to modify gravity to avoid postulating an undetected dark matter, the void solution to the SNe Ia supernovae observations does not require a modification of Einstein's gravity theory due to its claim that there is no ``dark energy''. An underdense void expands faster than its more dense surrounding galaxies, whereby younger supernovae inside the void would be observed to be receding more rapidly than older supernovae outside the void. By assuming the LFRW homogeneity, the mistaken conclusion would be reached that the expansion rate of the universe is accelerating, while in fact the local void and the global universe are actually decelerating. The local void scenario has been investigated by many authors. For a review, see~\\cite{Celerier} and for more recent citations see~\\cite{Sarkar}. There are other problems occurring in the void model, in addition to giving up the cosmological Copernican Principle. To fit the CMB data and other pertinent cosmological data, the radial size of the void required to fit the supernovae data is hundreds of Mpc to Gpc, and we must be close to the center of the void to avoid an unacceptably large CMB dipole~\\cite{Moffat,Alnes}.  Another problem to be considered in the inhomogeneous void model is the assumed initial conditions of the universe. Strictly speaking, simple models of inflation would be inconsistent with the existence of a large inhomogeneous void. However, there exists an alternative solution to the initial value horizon and flatness problems. This is based on a bimetric gravity variable speed of light (VSL) model~\\cite{Moffat3}. It has been demonstrated by Magueijo~\\cite{Magueijo} that this model in conjunction with the Dirac-Born-Infeld model can predict almost scale invariant primordial fluctuations in agreement with observations without inflation. This model would not necessarily negate the possibility of large, nonlinear and inhomogeneous voids.  In the following, we will review discoveries made in the early investigations of the void model~\\cite{Moffat,Moffat2}, and consider the age problem in cosmology and how the inhomogeneous cosmology can ameliorate this problem. We also consider the possible amplification of growth structure and its early onset in the inhomogeneous LTB model. We do not consider the important issue of confronting the void model with the Cosmic Microwave Background (CMB) WMAP data~\\cite{Komatsu}. Attempts to do this have been made in the recent literature~\\cite{Sarkar,Tomita2,Tomita3,Biswas,Zibin,Garcia}. These investigations remain inconclusive as to how well the void model can produce the acoustical power spectrum, the matter power spectrum associated with large scale surveys of galaxies and the observed baryon acoustical oscillations (BAO) in the correlation function. The role of structure growth in the inhomogeneous models plays an important role in these cosmological calculations, as does the importance of contributions from ``warm\" dark matter such as neutrinos with a mass not violating recent upper bounds obtained from WMAP observations. We conclude that the void model has not been ruled out. However, the best fits to the CMB WMAP data and the BAO appear to require a low value of $H_0\\sim 44\\,{\\rm km\\,s^{-1}\\,Mpc^{-1}}$ compared to measurements of $H_0$~\\cite{Jackson}. We should nonetheless be cautious about the true value of $H_0$, for it has a long history of changing its value with time. Moreover, in an inhomogeneous cosmology the value of the Hubble constant is position dependent. The local measured value of $H_0$ near the center of the void is expected to be larger than in distant parts of the universe outside the void. This could accommodate fits to the CMB data at the surface of last scattering with lower values of $H_0$ than would be measured by an observer close to the center of the void. ",
        "conclusions": "We have investigated whether an underdense, nearly spherically symmetric void can successfully describe cosmological observations. In spite of the impressive success of the standard $\\Lambda$CDM model, based on a homogeneous, isotropic and spatially flat LFRW universe, in fitting a large amount of data including both the type Ia supernovae data and the CMB data, there are problems with the standard model: the undetected dark matter, the cosmological constant problem and the anti-Copernican coincidence problem associated with the vacuum density dark energy. To avoid the generic problems with dark energy the possibility of having a cosmology with a large underdense void, embedded in an Einstein-de Sitter spacetime has been proposed. This avoids having to modify Einstein's gravity theory at least as far as the issue of dark energy is concerned. We have shown that already in an early publication~\\cite{Moffat2} the solution with a suitable void density profile, anticipated the supernovae observation of a dimming of the supernovae light, because the void expands faster than the surrounding overdense shell of matter. However, in the void model the acceleration of the universe, as interpreted in an LFRW universe, is a misinterpretation of the observations, for the void and the universe outside the void are in fact decelerating. This means that in the void model we can discard the idea of an undetected dark energy.  The void solution also has the ``philosophical\" complication of giving up the cosmological Copernican Principle by requiring that the observer be near the center of the void. Whether cosmology demands a fully Copernican Principle is still an observationally unresolved question. However, the void scenario can only be convincing, if it is fully consistent with the accurate CMB WMAP data, and the observed matter power spectrum based on large scale galaxy surveys, the ISW effect and the BAO data. We cannot rule out the possibility yet that the inhomogeneous void cosmology can fit this data satisfactorily. The question as to how perturbation growth of structure both in the primordial cosmology and in later time cosmology is affected by inhomogeneous perturbation calculations is not yet fully understood. This means that we have to be cautious about applying computer codes such as CMBFAST and CAMB, which were designed exclusively to give fits to data based on the standard $\\Lambda$CDM model, to the more general inhomogeneous cosmology models.  The problem of the age of the universe and the amplification of perturbation calculations of structure growth and correlation functions have been addressed here. We found that as already discovered 1n one of the first phenomenological investigations of the inhomogeneous void cosmology~\\cite{Moffat,Moffat2}, based on the exact pressureless LBT model, the age of the universe problem can be solved in these models and that an increased amplification and early onset of structure growth is to be expected.  An alternative inhomogeneous cosmology model based on the Szekeres-Szafron~\\cite{Szafron} exact solution of Einstein's field equations has been investigated~\\cite{Moffat9}. The solution does not in general have any symmetry and incorporates pressure as well as energy density. A particular cylindrically symmetric Szafron solution was studied, which can describe an early universe with non-uniform inflation, has homogeneous LFRW behavior at the surface of last scattering and can become inhomogeneous at a late time of the expansion of the universe.  {\\bf Acknowledgements:}  I thank Bruce Bassett and Viktor Toth for stimulating discussions and Viktor Toth for a numerical calculation. This research was supported by a grant from NSERC. The Perimeter Institute is supported in part by the Government of Canada through NSERC and by the Province of Ontario through MEDT."
    },
    "0910/0910.3725_arXiv.txt": {
        "abstract": "The disk corona evaporation model extensively developed for the interpretation of observational features of black hole X-ray binaries (BHXRBs) is applied to AGNs. Since the evaporation of gas in the disk can lead to its truncation for accretion rates less than a maximal evaporation rate, the model can naturally account for the soft spectrum in high luminosity AGNs and the hard spectrum in low luminosity AGNs. The existence of two different luminosity levels describing transitions from the soft to hard state and from the hard to soft state in BHXRBs, when applied to AGNs, suggests that AGNs can be in either spectral state within a range of luminosities. For example, at a viscosity parameter, $\\alpha$, equal to 0.3, the Eddington ratio from the hard to soft transition and from the soft to hard transition occurs at 0.027 and 0.005 respectively. The differing Eddington ratios result from the importance of Compton cooling in the latter transition, in which the cooling associated with soft photons emitted by the optically thick inner disk in the soft spectral state inhibits evaporation.  When the Eddington ratio of the AGN lies below the critical value corresponding to its evolutionary state, the disk is truncated. With decreasing Eddington ratios, the inner edge of the disk increases to greater distances from the black hole with a concomitant increase in the inner radius of the broad line region, $R_{BLR}$. The absence of an optically thick inner disk at low luminosities ($L$) gives rise to region in the $R_{BLR}-L$ plane for which the relation $R_{BLR} \\propto L^{1/2}$ inferred at high luminosities is excluded. As a result, a lower limit to the accretion rate is predicted for the observability of broad emission lines, if the broad line region is associated with an optically thick accretion disk.  Thus, true Seyfert 2 galaxies may exist at very low accretion rates/luminosities. The differences between BHXRBs and AGNs in the framework of the disk corona model are discussed and possible modifications to the model are briefly suggested. ",
        "introduction": "The structure of accretion disks surrounding compact objects is central to our understanding of both stellar mass black holes in X-ray binary systems and supermassive black holes in the nuclei of galaxies.  Among the issues requiring elucidation is the physical description of the accretion flow in the sources' spectral states. Toward this end, Shakura \\& Sunyeav (1973) pioneered the development of optically thick and geometrically thin accretion disks within the framework of a turbulent $\\alpha$ disk model, which provided an explanation of the thermal state of these sources.  Recognizing the possibility of physically distinct states in the disk to explain the hard X-ray state of Cyg X-1 (e.g., Agrawal et al. 1972; Tananbaum et al. 1972), Shapiro, Lightman, \\& Eardley (1976) and Ichimaru (1977) proposed a geometrically thick and optically thin accretion flow model. More recently, Narayan \\& Yi (1994, 1995a, 1995b) and Abramowicz et al. (1995) have stressed the importance of the advection of accretion energy stored in the accreting gas, leading to radiative inefficient accretion in the so-called advection-dominated accretion flow (ADAF).  The investigation of the features of these ADAF models and their application to the spectral states of black hole X-ray binary systems and AGNs has been discussed in the reviews of Narayan, Mahadevan, \\& Quataert (1998), Narayan (2005), and Narayan \\& McClintock (2008).  The high/soft state spectrum of BHXRBs is generally believed to originate from an accretion disk extending to the innermost stable circular orbit (ISCO) as developed by Shakura \\& Sunyaev (1973), whereas the low/hard state spectrum is thought to originate in a geometrically thick, optically thin, hot corona or ADAF in the immediate vicinity of the black hole (see Narayan \\& Yi 1994; 1995a, b). We note that the observational data of accreting black holes commonly point to the coexistence of hot and cool accretion flows, which is likely to be either in the form of an inner ADAF connecting to an outer geometrically thin disk or an ADAF-like corona lying above a standard thin disk extending to the ISCO. The former configuration has been studied by Esin et al. (1997) and applied to both quiescent BHXRBs and low-luminosity AGNs.  In contrast, the latter configuration was proposed to explain the power-law X-ray emission and lower-frequency blackbody component observed in both BHXRBs and AGNs (e.g., see Liang \\& Price 1977; Haardt \\& Maraschi 1991; Nakamura \\& Osaki 1993). In the context of these studies, only the disk evaporation model (e.g. Meyer et al. 2000a) has been investigated in detail to elucidate the nature of the interaction between the disk and the corona.  Such an evaporation model has been reasonably successful in providing an interpretive framework of the many features observed in BHXRBs. For example, the spectral transition between the low/hard and high/soft state (Meyer et al. 2000b), the disk truncation and formation of an ADAF in the low/hard state (Liu et al. 1999; Meyer-Hofmeister \\& Meyer 2003), the luminosity hysteresis between the hard-to-soft and soft-to-hard transitions (Meyer-Hofmeister et al.  2005; Liu et al. 2005), the intermediate state (Liu et al. 2006; Meyer et al. 2007), and the inner cool disk component in the low/hard state (Liu et al. 2007; Taam et al.  2008) can all be understood as consequences of the disk evaporation/condensation process. The model provides a physical basis for understanding the observational phenomenology, although details remain to be addressed.  As is well known, observations of AGNs exhibit many similar features to BHXRBs. Of importance is the existence of the fundamental plane (Merloni et al. 2003; Falcke et al. 2004; Wang et al. 2006) relating the radio luminosity, X-ray luminosity and black hole mass which extends from stellar masses to supermassive black holes.  In addition to these relations, the differing spectral shapes between high luminosity AGNs (HLAGNs) and low luminosity AGNs (LLAGNs) (e.g., Elvis et al. 1994; Ho 1999; Ho 2008 and reference therein) are widely observed.  Specifically, the HLAGNs, such as quasars and some Seyfert galaxies, display evidence of disk accretion as inferred from the presence of a ``big blue bump'' and hence are the analogues of high/soft state BHXRBs (Maccarone et al. 2003). On the other hand, LLAGNs exhibit hard X-ray emission as illustrated, for example, by our Galactic Center (Narayan \\& Yi 1994; Yuan et al.  2003; Narayan 2005) and/or evidence of radio jets and are analogues of the low/hard state BHXRBs (for a review see Ho 2008). Spectral features seen in narrow line Seyfert 1 galaxies (NLS1) are similar to BHXRBs at the very high states (e.g. Pounds et al. 1995; Pounds \\& Vaughan 2000). Although there exist differences (e.g., a soft X-ray excess is often seen in AGNs), a phenomenology based on a unification scheme for the BHXRBs and AGNs has been proposed (Falcke et al.  2004; Fender et al. 2004; Ho 2009).  In these models, the transition from an outer disk to an inner ADAF is one of the primary issues remaining to be clarified.  Among the mechanisms facilitating the state transition (see Honma 1996; Manmoto \\& Kato 2000; Meyer et al. 2000b; R\\`o\\.za\\`nska \\& Czerny 2000a,b; Spruit \\& Deufel 2002; Lu et al. 2004; Dullemond \\& Spruit 2005), the disk corona evaporation model provides a promising explanation for not only the state transition, but also the observational features of BHXRBs aforementioned.  It is, thus, natural to confront the disk corona evaporation model to observations of AGNs to determine its applicability to accreting black holes systems in general.  Investigation of the model reveals that the evaporation feature is independent of the black hole mass. In particular, the coronal temperatures in the inner region of disks around supermassive black holes and in BHXRBs are the same.  Furthermore, the radial distribution of the evaporation rate and the mass flow rate when scaled to the black hole mass for both AGNs and BHXRBs are also independent of the mass. Thus, the model can, in principle, be applied to AGNs (Liu \\& Meyer-Hofmeister 2001), especially given the analogues between the observational features of AGNs and BHXRBs.  In this regard, a spectral state transition for AGNs with a luminosity of a few percent of the Eddington value, the truncation of a thin disk for LLAGNs, the presence of iron lines related with both the disk and the corona, and the occurrence of radio jets associated with the evaporative process in disk coronal flows may be expected.  To investigate the disk evaporation/condensation model for AGNs, we outline the basic physics and assumptions of the theoretical model in \\S 2 and discuss its applicability as an interpretative framework for several AGN properties in \\S 3. The implications of the model and the issues that remain to be addressed due to differences between AGNs and BHXRBs are presented in \\S 4. ",
        "conclusions": "We have shown that a model developed for BHXRBs based on a disk-corona interaction provides a promising interpretative framework for explaining the characteristic observational features of AGNs. The evaporation of gas leads to truncation of the disk for accretion rates less than a maximal evaporation rate.  This characteristic naturally accounts for the distinct states of AGN, viz., a soft spectrum with big blue bump at high luminosities and a hard spectrum at low luminosities. Due to the existence of differing critical mass flow rates describing spectral transitions from the hard state to the soft state and the soft state to the hard state, AGNs characterized by the same Eddington ratio within the critical luminosity interval can be in either state.  Associating the characteristic scale of the broad line region, $R_{BLR}$, with the disk truncation radius for LLAGNs, the model predicts that the region excluding the presence of broad lines increases with decreasing luminosity for luminosities less than about $0.005 -  0.027 L_{Edd}$ for an $\\alpha$ viscosity parameter of 0.3. Hence, $R_{BLR}$ departs from the relation $R_{BLR} \\propto L^{1/2}$ at these luminosity levels. As a consequence, broad emission lines are only expected for AGNs characterized by mass accretion rates above a given value, if the broad line region is associated with an optically thick accretion disk. Thus, true Seyfert 2 galaxies are expected at sufficiently low luminosities.  As an additional consequence of the disk truncation, a thermal red bump is expected to be present at low Eddington ratios. This expectation is in qualitative agreement with the presence of the so called ``big red bump'' in LLAGNs (see Ho 1999; 2008). A detailed study on the disk truncation and a comparison with the observational properties of the big red bump is reserved for a future paper.  Although the model's basic features provide theoretical support for the hypothesis that the accretion process for supermassive black holes in AGNs is similar to the process for stellar mass black holes in BHXRBs, several observational differences exist between these two types of systems.  In the following, we discuss these differences and their possible implications for the model.  \\subsection{Strong coronal flow in AGNs}  As the evaporation is saturated to the maximal rate at high accretion rates (Eddington ratios), the accretion flow in the corona cannot increase with the accretion rate after a transition to the soft state. Thus, the accretion flow within the disk (rather than corona) increases with increasing accretion rate. With efficient Compton scattering, the corona component could be even weaker at high accretion rates because the increased disk radiation leads to strong Compton cooling, resulting in the condensation of coronal gas and the reduction of the optical depth in the corona. If we adopt the maximal evaporation rate without Compton cooling as an upper limit for the corona flow, we obtain an upper limit on the ratio of the mass inflow rates through the corona and disk as \\begin{equation} {\\dot m_c\\over \\dot m_d}={1\\over {\\dot m\\over 0.027}-1}, \\end{equation} which is also an upper limit to the ratio of the coronal luminosity to the disk luminosity of a source in the high/soft state.  For an accretion rate of 0.1 the coronal luminosity is no more than 37\\% of the disk luminosity and is less than $\\sim 3\\%$ of the disk luminosity if the system accretes at the Eddington rate. This implies that the spectral index between optical and X-rays, $\\alpha_{\\rm ox}$, is very large in luminous AGNs.  However, observations reveal that the X-ray luminosity is comparable with the optical luminosity in some HLAGNs, indicating a strong corona in AGNs (which is similar to the very high state in BHXRBs). To produce such strong X-ray radiation in luminous AGNs, the frictional heating (in the model) is insufficient. To address this conundrum, it has been suggested that a large fraction of accretion energy is transported from the disk to the corona, perhaps due to a magnetic buoyant instability (Liu et al. 2002b; 2003; Cao 2009).  Alternatively, the viscosity parameter $\\alpha$ in these HLAGNs may be large, supporting a strong corona by means of efficient evaporation.  Such a possibility could be a consequence of a larger magnetic Prandtl number in AGNs than in BHXRBs (see Lesur \\& Longaretti 2007).   \\subsection{Additional phenomenology}  In addition to their apparent strong coronae, AGNs exhibit other observational features which differ from BHXRBs. For example, among the luminous AGNs $\\sim 10\\%$ are radio-loud QSOs, while jets in BHXRBs appear to quench whenever the objects transit to the high/soft state. This difference may reflect that jets in the high state of BHXRBs are weak or the properties of the magnetic field in the corona, related to the magnetic Prandtl number, are modified (see above) to facilitate jet formation in AGNs.  We also note that soft X-ray excesses are seen in AGNs at high Eddington ratios, but not in BHXRBs. Such excesses appear to be correlated with the presence of the big blue bump.  The origin of this excess is unknown, but it has been interpreted in terms of Comptonization in the disk (Kawaguchi et al. 2001) or resulting from a depression of the continuum at X-ray energies in the range 0.7-5 keV by relativistically broadened line transitions of O VII and O VIII (Gierlinski \\& Done 2004). Its restriction to AGNs remains to be clarified.  The environment of a galactic central region can significantly differ from that for a stellar mass black hole in an X-ray binary system. For example, significant high temperature diffuse gas has been observed in the Galactic Center, with Sgr A$^*$ accreting at very low Eddington ratios, which may also be present in some LLAGNs. Such hot gas would affect the boundary condition on the disk. However, in the case of hot gas joining the evaporating gas in the corona, the corona is not changed significantly. This is a consequence of the fact that the evaporation of the disk is depressed by the directly inflowing gas due to the pressure and energy balance between the disk and the corona. If the density in the corona is too high, the efficient cooling to the corona can even lead to the coronal gas condensing to the disk. Consequently, the corona flow remains similar to the case of evaporation-fed corona.  In other words, for a given accretion rate, independent of whether the disk is fed at the outer boundary, or the disk (corona) is partially fed in the form of cold (hot) gas, the accretion flows through the disk and through the corona in regions of a few hundred Schwarzschild radii are expected to be similar.  Finally, the disk instability model may not give rise to large amplitude outbursts (and hence transitions between the low/hard state and the high/soft state) due to the ionization instability in AGNs as in BHXRBs because the $\\alpha$ parameter in the cold state may not significantly differ from its value in the hot state (see Menou \\& Quataert 2001). In this case the variability would be of small amplitude and state transitions would be less dramatic. AGNs would therefore primarily reflect the mass inflow rate from its surrounding region rather than a temporal variation in a non-steady disk. On the other hand, gravitational instabilities may operate in AGNs to give rise to more significant time dependent accretion.  In the future, we plan to carry out additional investigations to quantitatively confront the disk corona model with the detailed observations of AGNs, with comparisons based on a two-temperature corona model (Liu et al. 2002a) incorporating the effect of magnetic fields (Qian et al. 2007) and viscosity (Qiao \\& Liu 2009).  In addition, a study of their differences with respect to BHXRBs will be explored to examine whether they are one of degree rather than of kind."
    },
    "0910/0910.3987_arXiv.txt": {
        "abstract": "Positronium is the short-lived atom consisting of a bound electron-positron pair.  In the triplet state, when the spins of both particles are parallel, radiative recombination lines will be emitted prior to annihilation.  The existence of celestial positronium is revealed through gamma-ray observations of its annihilation products. These observations however have intrinsically low angular resolution. In this paper we examine the prospects for detecting the positronium recombination spectrum.  Such observations have the potential to reveal discrete sources of \\pos\\ for the first time and will allow the acuity of optical telescopes and instrumentation to be applied to observations of high energy phenomena.  We review the theory of the positronium recombination spectrum and provide formul\\ae\\ to calculate expected line strengths from the \\pos\\ production rate and for different conditions in the interstellar medium.  We estimate the positronium emission line strengths for several classes of Galactic and extragalactic sources.  These are compared to current observational limits and to current and future sensitivities of optical and infrared instrumentation. We find that observations of the \\psa\\ line should soon be possible due to recent advances in near-infrared spectroscopy. ",
        "introduction": "Positronium (Ps) is the short-lived atom consisting of a bound electron-positron pair.  Its existence was predicted by \\citet{moh34} shortly after the discovery of the positron by \\citet{and33}.  Furthermore, \\citet{moh34} suggested that Ps should be observable by its recombination spectrum and discussed its properties.  Ps was first observed by \\citet{deu51} in a laboratory experiment designed to measure the lifetimes of positrons in gases.  Ps is unstable and annihilates into two or more photons with a combined energy of 1022 keV, i.e. the rest mass of the atom (see \\S~\\ref{sec:ann}).  However, as discussed in \\S~\\ref{sec:phys}, under certain conditions Ps is stable for long enough that electronic transitions can occur between the quantum states of the system, giving rise to a recombination spectrum.  Since the reduced mass of Ps is half that of H, the wavelengths are twice those of the corresponding H series.  For example, Ps Lyman $\\alpha$ (\\psla) has a wavelength of 2431\\AA\\ and Ps Balmer $\\alpha$ (\\psa) has a wavelength of 1.313$\\mu$m.  \\psla\\ was first observed in laboratory experiments by \\citet{can75}.  So far, however, no detections have been made of recombination lines from celestial Ps.  Nevertheless the existence of celestial Ps is inferred from observations of the 511keV annihilation line and continuum annihilation radiation seen both from the Sun (\\citealt{chu73}) and from the Galactic centre (\\citealt{lev78}), combined with the fact that in most astrophysical environments \\pos\\ are expected to form Ps before annihilation (cf.\\ \\S~\\ref{sec:phys}).  The successful observation of Ps recombination lines would not only be of interest in its own right, but would also have important consequences for many areas of astrophysics.  We now describe the motivation behind a search for Ps recombination lines.  Current knowledge of naturally occurring antimatter from astrophysical sources is limited to the poor spatial resolution ($\\sim 3^{\\circ}$) of $\\gamma$-ray satellites.  The formation of Ps, however, provides a natural heterodyne process that converts the $\\gamma$-ray energies associated with the production of \\el-\\pos\\ pairs, to the optical/ near-infrared energies of Ps recombination. Thus observation of Ps recombination lines would allow the superior spatial and spectral resolution of optical/ near-infrared technologies to be applied to high energy astrophysical phenomena.  As a specific example, Ps observations could provide an important physical understanding of jets.  It is still uncertain whether high energy radio jets consist of an \\el-p plasma or \\el-\\pos\\ or some combination of both.  Most attempts to answer this question have had to rely on indirect methods of searching for the presence of \\pos, such as arguments based on the total energy budget of the jets (e.g.\\ \\citealt{cel93}; \\citealt{rey96}; \\citealt{war98}; \\citealt{hir05}), with inconclusive results.  Direct searches for the 511keV annihilation signature  in jets (\\citealt{mar07}) have likewise failed to provide conclusive results due to the poor point source sensitivity of $\\gamma$-ray telescopes.  If the positrons from the jet thermalise when the jet impacts the surrounding medium allowing the formation of Ps, and if the acuity of optical observations can be brought to bear on this problem to allow a successful detection of Ps recombination lines, this would have significant implications for our understanding of the physical processes responsible for the production of jets -- processes which take place deep within the AGN (\\citealt{bla95}).  The origin of the positrons responsible for the annihilation radiation observed from the Galactic centre is uncertain.  The current state-of-the-art observations are from the ESA's INTEGRAL $\\gamma$-ray observatory, which has detected extended Ps annihilation radiation centred on the Galactic centre with a FWHM of $\\approx 8^{\\circ}$ (\\citealt{kno05}; \\citealt{wei08b}).    The results are most consistent with a source of \\pos\\ originating from old stellar populations such as type Ia supernov\\ae , classical nov\\ae , and low mass X-ray binaries, although more exotic sources, such as dark matter annihilation, cannot be ruled out.  The angular resolution of INTEGRAL SPI is $\\sim 3^{\\circ}$, and is therefore rather insensitive to the detection of point sources.   The detection of Ps recombination lines would allow discrete point sources of \\pos\\ to be searched for on sub-arcsecond scales (see \\S~\\ref{sec:sources}) allowing direct  confirmation of the \\pos\\ production of specific sources for the first time.     Despite the benefits of detecting Ps recombination lines there have been very few attempts to do so, and there has been little work on the expected astrophysical signatures.   \\citet{mcc84} reviews the possibility of the optical detection of Ps, in particular the \\psla\\ line.  The conclusion reached is that a visual extinction of 1 or 2 magnitudes is enough to make the Balmer and Paschen lines more intense than the intrinsically more luminous Lyman lines.  Since most previous studies have focussed on the Galactic centre as  the most promising source of Ps recombination line emission (so far the only source other than the Sun in which Ps annihilation has been observed) and that at the Galactic centre the visual extinction is very high, it is not surprising that this idea has remained largely forgotten.  \\citet{burd97} raised the possibility of optical/ IR detection of Ps again, and made predictions of the expected fluxes of the most important lines.  However, the only published attempt to detect Ps recombination lines was made by \\citet{pux96} who searched for Ps Paschen-$\\beta$ from the Galactic centre in the $K$ band, and whose non-detection provides the only direct  experimental constraints on the emission.  Until recently, the prospect of a near-infrared (NIR) detection of the Balmer series of Ps was rather bleak due to the relative insensitivity of  NIR detectors and the strong foreground emission from the atmosphere.  However, we argue that the time is now right for a reappraisal of this important window into the high energy universe.  This renewed interest is motivated by several recent and imminent advances in technology.  NIR detectors have improved remarkably and are now capable of dark currents as low as 1\\el\\ per 1000 s, and read noises of $\\approx 8$\\el (\\citealt{smi06}).  Adaptive optics at NIR wavelengths can now deliver high Strehl ratios improving the signal to noise for faint sources.  The future generation of ELTs will deliver AO corrected images of fields of view of tens of arcminutes with huge gains in sensitivity from the vast collecting areas.  The JWST will provide unprecedented sensitivity at NIR wavelengths (\\citealt{gar06}).  Recent breakthroughs in photonic technologies promise superb background suppression in the NIR in the near future (\\citealt{bland04}; \\citealt{ell08}).  The purpose of this paper is to assess the probability of detecting Ps recombination spectra in light of these advances and to identify the most fruitful strategies  and targets for an observational campaign.  We collect together sources from a broad literature published over many decades and combine these to calculate the expected astrophysical signatures of Ps. The next section reviews the pertinent physical properties of Ps.  We then discuss possible sources of Ps recombination line emission, both Galactic and extragalactic, including predictions on the strengths of these sources.  Section~\\ref{sec:expt} then assesses the prospects of detection and observational strategies.  Finally we summarise our findings in section~\\ref{sec:summ}. ",
        "conclusions": "\\begin{tabular}{llllll} Source & \\multicolumn{2}{c}{Predicted Flux} & \\multicolumn{2}{c}{Upper limit}\\\\ & \\multicolumn{4}{c}{ph s$^{-1}$ m$^{-2}$ $\\mu$m$^{-1}$} \\\\ & \\psla\\ & \\psa\\ & \\psla\\ & \\psa\\ \\\\ \\hline Kepler SNR & 1.6 & 7.9 & 12$^{a}$ & 60$^{a}$ \\\\ LMXB GX 5-1 & 0.065 & 3.7 & 0.65$^{a}$ & 37$^{a}$ \\\\ Crab Pulsar & 130 & 110 & 190$^{a}$ & 160$^{a}$ \\\\ 3C 120 & 8.9 & 1.5 & 170$^{b}$ & 50$^{b}$ \\\\ \\end{tabular} \\end{center} \\end{table}  \\subsection{Expected sensitivity of current and future instruments} \\label{sec:expected}  We now examine the sensitivity of current and future instruments to Ps emission line radiation.  The signal to noise ratio is given by \\begin{equation} \\label{eqn:snr} SNR=\\frac{f_{\\lambda} t \\epsilon}{\\sqrt{f_{\\lambda} t \\epsilon+ B t \\epsilon + D t + R^{2}}} \\end{equation} where $t$ is the exposure time, $\\epsilon$ is a factor which takes into account the efficiency of the instrument and collecting area of the telescope, $B$ is the background, $D$ is the detector dark current and $R$ is detector read out noise. Since the expected Ps recombination line emission is faint in all cases it is likely that the background will also include continuum emission from the source in question as well as the usual night sky background etc.  We first address the issue of whether it is more efficient to observe \\psla\\ or \\psa.  \\psla\\ is always intrinsically brighter than \\psa.  It does not follow however, that it is more efficient to observe \\psla\\ than \\psa.  This is because the interstellar extinction, the line width, the background and the seeing all very with wavelength.  Interstellar extinction is greater for \\psla\\ than for \\psa.  This means that the integrated line strength for \\psa\\ is greater than that for \\psla\\ for visual extinctions $A_{\\rm V} \\gsim 0.65$ mag, where the exact value depends on the temperature of the ISM in which the Ps is forming, see Figure~\\ref{fig:extinct_total}.  However, the line width is also greater for \\psa\\ compared to \\psla\\ so the line flux density, $f_{\\lambda}$, which is more meaningful to compare to the background is greater for \\psa\\ for visual extinctions $A_{\\rm V} \\gsim 1.2$ mag, see Figure~\\ref{fig:extinct_density}.  \\begin{figure} \\centering \\includegraphics[scale=1.0]{avtotal.eps} \\caption{The ratio of \\psa/\\psla\\ for integrated line strengths as a function of visual extinction for different temperatures in the ISM.  The horizontal dot-dashed line marks the point at which the integrate line strength of both lines are equal.  Above the dashed line \\psa\\ is brighter than \\psla.} \\label{fig:extinct_total} \\end{figure}  \\begin{figure} \\centering \\includegraphics[scale=1.0]{avdensity.eps} \\caption{The ratio of \\psa/\\psla\\ for line flux densities as a function of visual extinction for different temperatures in the ISM.  The horizontal dot-dashed line marks the point at which the integrate line strength of both lines are equal.  Above the dashed line \\psa\\ is brighter than \\psla.} \\label{fig:extinct_density} \\end{figure}   A proper assessment of the most fruitful observing strategy depends on other factors, most notably the background emission and the efficiency of the instruments.  Therefore we now discuss the different backgrounds for \\psla\\ and \\psa, and then determine the expected sensitivity of observations compared to the expected brightnesses of example objects.  The U band background at a typical good observing site is $U=21.5$ mag arcsec$^{-2}$, which is equivalent to a background of $B_{\\rm U}=140$ ph s$^{-1}$ m$^{-2}$ $\\mu$m$^{-1}$ arcsec$^{-2}$, which we assume is also appropriate for observations at 2430\\AA.  The background in the near-infrared is dominated by atmospheric hydroxyl emission lines.  Thus the background depends sensitively on the wavelength and the resolution of the observations.  Figure~\\ref{fig:nirback} shows the logged background around the \\psa\\ line, taken from the model of \\citet{ell08}.  There is a very bright OH emission line very close to the \\psa\\ line.  Note that Figure~\\ref{fig:nirback} is at the intrinsic resolution of the night sky, i.e.\\ the width of the lines is the natural width due Doppler and pressure broadening; when observed with a spectrograph, the resolution and scattering of the spectrograph will cause the OH line to severely contaminate the \\psa\\ line.   Subtracting the OH line results in a high Poisson and systematic noise (see \\citealt{dav07}; \\citealt{ell08}). Indeed, we postulate that this is the reason that \\psa\\ has never been observed serendipitously.  \\begin{figure} \\centering \\includegraphics[scale=0.5]{nirback.eps} \\caption{The background around the \\psa\\ line is shown by the thin grey line.  The flux is logged to emphasise the fainter lines and the interline continuum.  The dashed line marks the position of the \\psa\\ line.  The thick black line shows the background spectrum after OH suppression.} \\label{fig:nirback} \\end{figure}  The brightness of the near-infrared sky is problematic for many areas of observational astronomy, as exemplified by the particular science case of detecting \\psa.  For this reason there have been significant technological developments aimed at tackling this problem.  The implementation of several of these projects is now imminent, promising much darker near-infrared skies in the near future.  These technological developments are the motivation behind our renewed interest in the long standing observational challenge of detecting Ps recombination emission lines.  There are two technologies in particular which should be very beneficial for the detection of \\psa.  These are the James Webb Space Telescope\\footnote{JWST http://www.jwst.nasa.gov/} (see e.g.\\ \\citealt{gar06}) and OH suppressing fibre Bragg gratings (\\citealt{bland04,bland08}).  The James Webb Telescope mission is designed specifically to achieve very low infrared backgrounds, both for imaging and moderate resolution spectroscopy  ($R\\approx 1000$--3000).  \\citet{bland04} and \\citet{bland08} present a novel photonic solution to the problem of the NIR night sky for ground based telesopes.  Used in conjunction with adaptive optics and large telescopes OH suppression offers very low backgrounds; fibre Bragg gratings can suppress the OH lines by factors of $\\approx 30$dB at a resolution 10,000.  For an estimate of the performance and impact of FBG OH suppression on astrophsyical observations see \\citet{ell08}; for a demonstration of their performance in on-sky tests see \\citet{bland09}.  Thus future NIR spectrographs will benefit from much lower backgrounds.  The thick black line in Figure~\\ref{fig:nirback} shows the expected background around the \\psa\\ line after suppression with FBG technology based on current performance. We note that at NIR wavelengths the background can be further reduced with adaptive optics which permits the use of a smaller aperture to collect the same object flux. After OH suppression the background flux density  at the wavelength of \\psa\\ is less than that of \\psla, assuming the interline continuum in the model of \\citet{ell08} is correct.  The calculation of the difference in backgrounds is complicated by the fact that the line width of \\psa\\ is greater than that of \\psla\\ and so the background must be integrated over a larger wavelength range for \\psa.  Conversely the seeing is smaller at longer wavelengths, and so the \\psa\\ background can be measured from a smaller aperture than for \\psla; this effect is even more pronounced with adaptive optics which can further reduce the necessary aperture in the infrared, but not in the ultraviolet.  In order to take into account all these various effects we have calculated the \\psla\\ and \\psa\\ flux limits to obtain a signal-to-noise of 10 as a function of time for backgrounds corresponding to \\psla\\ in natural seeing and \\psa\\ in natural seeing and adaptive optics corrected and with and without OH suppression.  The results are shown in Figure~\\ref{fig:fluxlimits}.  (Note that here we quote the results as total line strengths since the signal to noise calculation requires integrating the emission line and the background over the line width.)  We assume that the Ps lines are resolved, and the backgrounds integrated over the line width are 0.0272, 13.05 and 0.197 ph s$^{-1}$ m$^{-2}$ arcsec$^{-2}$ for \\psla, \\psa\\ and OH suppressed \\psa\\ respectively.  We assume that natural seeing is 1 arcsec in the UV and 0.5 arcsec in the NIR and that AO delivers a PSF of 0.1 arcsec FWHM.  We assumed a 8m telescope and a total system efficiency of 0.2.  \\begin{figure} \\centering \\includegraphics[scale=1]{fluxlimits.eps} \\caption{The flux limits to obtain a signal to noise of 10 as a function of time for \\psla\\ and \\psa\\ under various assumptions of the background (see text for details).} \\label{fig:fluxlimits} \\end{figure}  The limiting fluxes are compared to the average expected flux for microquasars (specifically we compare to GROJ1655-40 the brightest confirmed mis-aligned microquasar, cf.\\ Table~\\ref{tab:mqso}) as calculated above in \\S~\\ref{sec:mqso}.  Both the \\psla\\ and the \\psa\\ lines should easily be detectable.    We also compare to the fluxes for 3C 120, assuming the same backgrounds as for Ps at $z=0$ (note that the NIR background is a strong function of wavelength so this assumption is approximate in the case of no OH suppression).  In this case the Ps emission lines are too faint to be detected in natural seeing and natural backgrounds for both \\psla\\ and \\psa.   However with OH suppression and or AO correction \\psa\\ should be detectable.   \\subsection{Simulations}  We have estimated the Ps recombination emission line strengths for a wide variety of sources; the expected fluxes are fainter than current observational limits for \\psla\\ and marginal for \\psa.  However with the development of OH suppressing technologies  and the future launch of the JWST (\\citealt{gar06}) we are on the threshold of a new era in NIR spectroscopy with much lower backgrounds and consequently much more sensitivity.  The improved sensitivity will make the detection of \\psa\\ lines feasible.  The most promising targets for observation are mis-aligned microquasars and active galactic nuclei.  Emission from other candidates is generally too diffuse.  Mis-aligned microquasars should have a spot of Ps emission where the jet impacts the secondary star.  AGN should have similar spots of emission where the jet impacts gas clouds surrounding the nucleus.  In the case of AGN there will be bright emission from the nucleus itself, which must be subtracted off to reveal the faint Ps emission.  We have simulated observations of \\psa\\ from an AGN.  We use the observations of NGC~4151 by \\citet{sto09} as our template for the the AGN emission.  Specifically we model the spectrum of the region 0.7 arcsec E of the nuclues (see their figure 2, panel B), taking emission line fluxes from their table 1 and using a continuum value of $6 \\times 10^{-16}$ erg s$^{-1}$ cm$^{-2}$ \\AA$^{-1}$.  We model the \\psa\\ line by taking the flux for 3C120 based on the calculations of \\citet{mar07} (see Table~\\ref{tab:source_summary}) and scaling it to the redshift of NGC~4151.  The background spectrum is taken from \\citet{ell08}.  The simulated components of the spectrum are shown in Figure~\\ref{fig:ngc4151sim}.  \\begin{figure} \\centering \\includegraphics[scale=1.0]{n4151sim.eps} \\caption{Comparison of the components of the simulated spectrum of \\psa\\ emission from an AGN.  The emission from the AGN is shownby the thick black line, and is based on observations of NGC~4151 by \\protect\\citet{sto09}.  The \\psa\\ emission line is shown by the dashed black line and is based on calculations of the \\pos\\ production of 3C120 by \\protect\\citet{mar07}.  The near-infrared background is shown by the thin black line and is based on the models of \\protect\\citet{ell08}.} \\label{fig:ngc4151sim} \\end{figure}  The simulations assume OH suppressed observations on an 8m telescope with a 1 arcsec diameter fibre bundle as described in \\citet{ell08}, with adaptive optics correction delivering a Strehl ratio of 0.3.  The spectral resolution was assumed to be  $R=3000$.  The simulations assume a varying night-sky  background and systematic errors in wavelength calibrations between reads.  For further details see \\citet{ell08}.  Since the AGN continuum is much brighter than the \\psa\\ it must be treated as background, and observations will be required of the AGN plus \\psa\\ and of the AGN alone.  For this reason we anticipate that integral field spectroscopy will allow the greatest chance of a successful detection, since fluxes from many regions of the AGN can be  differenced, nullifying the requirement to know \\emph{a priori} in which region Ps might be forming.  For the purposes of the simulation we assume that the AGN continuum can be precisely scaled between these two locations (i.e.\\ in the simulations we assume the AGN emission does not change).  The simulation assumes a 6hr observation (composed of 12$\\times$30 min individual exposures) of both regions.  The resulting background and continuum subtracted spectrum is shown in Figure~\\ref{fig:ngc4151psa}.  The \\psa\\ line is clearly visible.  Without OH suppression the \\psa\\ is much fainter than the noise from the residual sky lines, and any detections would be marginal.  \\begin{figure} \\centering \\includegraphics[scale=0.6]{ngc4151psa.eps} \\caption{A background and continuum subtracted simulated observation of \\psa\\ from an AGN.   The simulations are shown with OH suppression by the thick black line, and without by the thin grey line.  The dashed line shows the expected location of the \\psa\\ line.  The line is clearly detected in the OH suppressed observation.} \\label{fig:ngc4151psa} \\end{figure}"
    },
    "0910/0910.1385_arXiv.txt": {
        "abstract": "While it is well known that neutrinos are emitted from standard core collapse protoneutron star supernovae, less attention has been focused on neutrinos from accretion disks.  These disks occur in some supernovae (i.e. \"collapsars\") as well as in compact object mergers, and they emit neutrinos with similar properties to those from protoneutron star supernovae.  These disks and their neutrinos play an important role in our understanding of gamma ray bursts as well as the nucleosynthesis they produce.  We study a disk that forms in the merger of a black hole and a neutron star and examine the neutrino fluxes, luminosities and neutrino surfaces for the disk.  We also estimate the number of events that would be registered in current and proposed supernova neutrino detectors if such an event were to occur in the Galaxy. ",
        "introduction": "Black hole - neutron (BH-NS) star mergers are potential progenitors of short duration gamma ray bursts (GRBs) and have been speculated to be the site of interesting nucleosynthesis. The neutrino emission from the accretion disk produced in BH-NS mergers plays an important role in each of these scenarios. One possible explanation for the energetic GRBs suggests as engine a black hole (BH) of several solar masses accreting matter from a disk. Simulations of compact object mergers have shown the formation of such disks \\cite{MacFadyen1999, Taniguchi2005, Lee1999, Rosswog}.  Neutrino transport, annihilation  and losses in the disk material would determine the GRBs production \\cite{Matteo02,Popham1999, Setiawan04}. Several studies have concluded that high accretion rates would provide the necessary conditions for triggering GRBs \\cite{Popham1999,SetiawanBHmerger, RuffertGRB-BH}. Another interesting aspect is the resulting nuclear products from the accretion disk around black holes (AD-BH). Lattimer and Schramm \\cite{Lattimer1974, Lattimer1976} speculated the possibility of a $r$-process in the accretion disk resulting of compact object mergers.  Material from inner crust of the merging neutron star (NS), with low proton fraction,  can be ejected from tidal tails giving place to a least a weak $r$-process.  Furthermore, hot accretion disk winds can also produce an r-process, due to the neutrino interactions \\cite{Surmanrprocess}.  Given the conditions of high temperature and density of AD-BH, we expect a copious amount of neutrinos  in the range of 10s of  MeV to be emitted. There are several detectors, both in operation and proposed, that could register neutrinos in this energy range, including  Super-Kamiokande (SK) \\cite{SK}, AMANDA \\cite{amanda,icecube}, KamLand \\cite{kamland}, ICARUS \\cite{icarus}, Ice Cube\\cite{amanda,icecube}, LANNDD\\cite{lanndd}, and HALO\\cite{halo}.  These detectors have been studied extensively for their ability to see a Milky Way supernova signal.  As detection of the next Galactic core collapse supernova is now within reach experimentally, we can begin to speculate on future detections of even rarer events such as BH-NS mergers and neutron star - neutron star (NS-NS) mergers. Roughly, NS-NS mergers are three orders of magnitude more rare than core collapse supernovae, and BH-NS mergers perhaps another order of magnitude still.  Given the rarity of these events, direct detection on a time scale of years would require a very large detector as we will discuss.  Nevertheless it is in interesting to examine the energy and strength of a signal that would originate from a BH-NS merger and compare with the signal of neutrinos emitted from a proto neutron star (PNS) at the center of  core collapse supernova.   Previous investigations of MeV scale signals in terrestrial detectors of neutrinos which originate from black hole accretion disks have focussed on disks that might form from the core collapse of a rotating massive star. Nagataki et al, investigated the luminosity, spectrum and counts at SK, for neutrinos from a collapsar \\cite{Nagatakicounts}. These authors used an analytical shape for the disk and from it derived the neutrino spectrum. Their results predict at least one neutrino event measured in the proposed 2Mega-ton water Cherenkov  detector, TITAND \\cite{Titand}, originating from  an AD-BH at 3 Mpc, when the total accretion mass, the initial mass, and the mass accretion rate are set to $30 M_\\odot$, $3M_\\odot$, and $0.1M_\\odot$s$^{-1}$, respectively. McLaughlin and Surman have considered the possible distinction of neutrino spectra when neutrinos originated in AD-BH versus a PNS \\cite{McLaughlin07}.  The AD-BH signal was found to have comparable energy spectra, but the disks primarily emit electron neutrinos and electron antineutrinos, which in addition to the timing of the signal could produce a unique signature after neutrino flavor transformation has been taken into account.  Determining the neutrino signal from a BH-NS merger is more complex.  In addition to calculating neutrino emission surfaces and estimating the effects of neutrino flavor transformation, general relativistic corrections are more important and one must determine world lines for the neutrinos which originate from different parts of the disk. In this paper we make estimates of neutrino events registered in several detectors for an AD-BH. We use a hyperaccreting disk model provided by Ruffert and Janka \\cite{Surmanrprocess,Ruffert2001}. We also calculate neutrino luminosities and fluxes for the disk. We take into account general relativity corrections for the neutrino energy and disk size. We discuss the consequences of neutrino oscillations in the fluxes and in our event counts. This paper is organized as follows: in section \\ref{disk model} we introduce the disk model, in section \\ref{neutrino surfaces} we present the reactions included to calculate neutrino surfaces. In section \\ref{Flux, Luminosity and Energy} we discuss our results for neutrino fluxes, energies and luminosities. In section \\ref{Neutrino Counts} we consider events rates in different detectors, in section \\ref{Neutrino Mixing} we take into account neutrino mixing, and in section \\ref{Conclusions} we discuss our conclusions. ",
        "conclusions": "\\label{Conclusions} The physics of NS-BH mergers is very important for our understanding of GRBs and the production of rare isotopes.  Using a state-of-the-art model for a NS-BH merger we have calculated neutrino surfaces, fluxes luminosities and energy averages for the neutrinos emitted from the emerging AD. Neutrino surfaces as seen from the $z$ axis have an irregular shape. Due to large cross sections the surface for $\\nu_e$ extends higher above the disk plane compared to the other neutrino flavor surfaces. We find that the redshifted neutrino energies are higher than in the case of a PNS. Our estimates are $E_{\\nu_e} \\approx 17.3 < E_{\\bar{\\nu_e}}  \\approx 23.4<E_{\\nu_x} \\approx 26$ MeV, which include redshift due to the strong BH gravitational field. The AD-BH neutrino luminosity from the NS-BH merger is also larger than the one of a PNS. We find a total neutrino luminosity $\\sim 10^{54}$ ergs/s; this is two orders of magnitude larger than the luminosity of a PNS. According to previous works this luminosity could be enough to trigger short GRB \\cite{SetiawanBHmerger}.  These results are based on a snapshot of the disk evolution. We have calculated the same quantities for a steady state model and found that the average energies are lower compared to the hydrodynamical model. This is due to the contribution of low temperature neutrinos emitted at large radial distances, $\\rho$, from the disk to the spectra in the steady state cases. We explored the influence of the spin parameter in the steady state disk using two different values.  We have found that higher spin parameters lead to higher neutrino energies and luminosities. As a consequence the neutrino counts per second are larger for the rapid rotating disk. However, if the fraction of binding energy released in neutrinos were the same for both, rotating and non-rotating disks, the signal in the rotating disk will be shorter resulting in less total counts. The order of magnitude of the luminosity for the steady state disk with $a$=0.95 is the same as for the 3D model and as it has been discussed could be a good candidate for triggering GRB \\cite{SetiawanBHmerger}.  As part of our estimates, we have explore the consequences of neutrino mixing on the neutrino energies and events registered. For this purpose, we consider a total of 8 oscillation scenarios, which include the normal and inverted hierarchy.  This demonstrates the range of possible fluxes and energies that can be expected for different flavors. These estimates of the luminosity and energies of neutrinos from the disks of BH-NS mergers have a number of applications.  Based in our results for the neutrino surface temperatures, we estimate the number of events that would be registered in several detectors, both proposed and in operation if a  BH-NS were to occur in the galaxy.  While such events are expected to be rare in comparison with the standard core collapse supernova, it is still important to understand the potential signal.  Our estimates can be used to rule out this object as the origin of a future neutrino signal, as well as to guide investigations of neutrino detectors for the distant future. For these estimates we have included corrections due to the BH gravitational field such as energy redshift and bending of the neutrino trajectories. In general we find that there will be more counts coming from an AD-BH that is formed in a neutron star - black hole merger than counts registered from a PNS. For example in SK we will register 1000 more counts if a NS-BH merger happens than if a supernova event occurs. However, the neutrino signal from AD-BH would last only 0.15 s compared to the 10 s estimate for a PNS. We have then a signal that is 100 times more luminous  and 100 times shorter than a PNS.  We warn that our estimated counts are sensitive to the time the signal lasts. And our estimate for the duration depends on the efficiency of converting gravitational energy into neutrino energy. Therefore, the total counts could change as much as a factor of two. Nevertheless, our results for neutrino energies, luminosity and fluxes are independent of this uncertainty. The amount of AD-BH neutrinos that could be detected in large scale facilities and the possible distinction of this signal from a PNS opens the question: Is there a way to distinguish the neutrinos from a BH-NS merger and a supernova? The timing of the signal and the energies of the neutrinos will be different, so this is an important clue.  Therefore, it is not likely that we will mistake neutrinos from a merger event as those from a core collapse supernova.  How could we guarantee a detection of BH-NS neutrinos? So far with a large scale facility like UNO we could detect neutrinos from AD-BH as far 4 Mpc, which covers galaxies in the local group such as Andromeda. To detect a BH-NS merger in neutrinos and compete with a gravitational wave detector such as Advance LIGO would take a three orders of magnitude larger detector."
    },
    "0910/0910.5665_arXiv.txt": {
        "abstract": "We present the stellar mass-size relations for elliptical, lenticular, and spiral galaxies in the field and cluster environments using {\\em HST}/ACS imaging and data from the Space Telescope A901/2 Galaxy Evolution Survey (STAGES). We use a large sample of $\\sim1200$ field and cluster galaxies, and a sub-sample of cluster core galaxies, and quantify the significance of any putative environmental dependence on the stellar mass-size relation. For elliptical, lenticular, and high-mass (${\\rm log}M_*/{\\rm M_{\\odot}}>10$) spiral galaxies we find no evidence to suggest any such environmental dependence, implying that internal drivers are governing their size evolution. For intermediate/low-mass spirals (${\\rm log}M_*/{\\rm M_{\\odot}}<10$) we find evidence, significant at the $2\\sigma$ level, for a possible environmental dependence on galaxy sizes: the mean effective radius $\\overline{a}_e$ for lower-mass spirals is $\\sim15$-$20$ per cent larger in the field than in the cluster. This is due to a population of low-mass large-$a_e$ field spirals that are largely absent from the cluster environments. These large-$a_e$ field spirals contain extended stellar discs not present in their cluster counterparts. This suggests the fragile extended stellar discs of these spiral galaxies may not survive the environmental conditions in the cluster. Our results suggest that internal physical processes are the main drivers governing the size evolution of galaxies, with the environment possibly playing a role affecting only the discs of intermediate/low-mass spirals. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.0871_arXiv.txt": {
        "abstract": "Star clusters are found in all sorts of environments, and their formation and evolution is inextricably linked to the star formation process. Their eventual destruction can result from a number of factors at different times, but the process can be investigated as a whole through the study of the cluster age distribution. Observations of populous cluster samples reveal a distribution following a power law of index approximately $-1$. In this work we use M33 as a test case to examine the age distribution of an archetypal cluster population and show that it is in fact the evolving shape of the mass detection limit that defines this trend. That is to say, any \\textit{magnitude-limited} sample will appear to follow a $dN/d\\tau=\\tau^{-1}$, while cutting the sample according to mass gives rise to a composite structure, perhaps implying a dependence of the cluster disruption process on mass. In the context of this framework, we examine different models of cluster disruption from both theoretical and observational standpoints. ",
        "introduction": "Star clusters are commonly used to trace the stellar content and star formation history of their host systems. The main limitations to this approach are the finite lifetime of a cluster (disruption) and evolutionary fading. From a theoretical point of view, the lifetime of a cluster should depend upon its initial mass and the properties of the environment in which it evolves (\\cite[Spitzer 1958]{spitzer58}\\footnote{This work showed that the lifetime due to encounters with interstellar clouds depends on the cluster density. Observations of young clusters have shown the radius to scatter tightly about 3.5~pc and can thus be considered a constant for any practical formulation. This immediately establishes the cluster mass as the main deciding factor of the lifetime of a cluster.}; \\cite[Baumgardt \\& Makino 2003]{bm03}). Observations have, however, produced \\textit{two} empirical disruption laws: \\begin{itemize} \\item Mass dependent disruption (MDD, \\cite[Boutloukos \\& Lamers 2003]{bl03}) \\item Mass independent disruption (MID, \\cite[Fall, Chandar \\& Whitmore 2005]{fcw05}) \\end{itemize}  In the MID model, `survivors' are selected on a purely random basis and a constant fraction is destroyed every age dex. Intriguing as it may be, this model clashes with several principles of cluster dynamics that would need to be revised considerably to accommodate it. In this contribution we test cluster dissolution and attempt to disentangle it from incompleteness and the statistical biases that have in the past limited such studies. We then compare theory to observations of the M33 cluster system. ",
        "conclusions": "\\label{sec:summary} Sample incompleteness causes a $\\tau^{-1}$ decline in the observed age distribution of cluster populations. Size-of-sample effects and thorough statistical methods can help to interpret the disruption process at play. The fundamental difference between magnitude- and mass-limited samples can then be used to discover the underlying physics. Our results from the nearby population of M33 are inconsistent with mass-independent disruption of a fixed fraction of clusters per age dex."
    },
    "0910/0910.5723_arXiv.txt": {
        "abstract": "Emission lines from metals offer one of the most promising ways to detect the elusive warm-hot intergalactic medium (WHIM; $10^5\\,\\K <T<10^7\\,\\K$), which is thought to contain a substantial fraction of the baryons in the low-redshift Universe. We present predictions for the soft X-ray line emission from the WHIM using a subset of cosmological simulations from the OverWhelmingly Large Simulations (\\owls) project. We use the \\owls\\ models to test the dependence of the predicted emission on a range of physical prescriptions, such as cosmology, gas cooling and feedback from star formation and accreting black holes. Provided that metal-line cooling is taken into account, the models give surprisingly similar results, indicating that the predictions are robust. Soft X-ray lines trace the hotter part of the WHIM ($T\\ga 10^6\\,\\K$). We find that the \\oviii\\ 18.97 \\AA\\ is the strongest emission line, with a predicted maximum surface brightness of $\\sim 10^2\\,$\\phot, but a number of other lines are only slightly weaker. All lines show a strong correlation between the intensity of the observed flux and the density and metallicity of the gas responsible for the emission. On the other hand, the potentially detectable emission consistently corresponds to the temperature at which the emissivity of the electronic transition peaks. The emission traces neither the baryonic nor the metal mass. In particular, the emission that is potentially detectable with proposed missions, traces overdense ($\\rho \\ga 10^2\\,\\rho_{\\rm mean}$) and metal-rich ($Z\\ga 10^{-1}\\,Z_{\\odot}$) gas in and around galaxies and groups. While soft X-ray line emission is therefore not a promising route to close the baryon budget, it does offer the exciting possibility to image the gas accreting onto and flowing out of galaxies. ",
        "introduction": "\\label{intro}  Observations of the low-redshift Universe are unable to account for a large fraction of the baryons predicted by cosmic microwave background (CMB) and nucleosynthesis measurements (e.g.\\ \\citealt{persic1992}; \\citealt{fukugita1998}; \\citealt{fukugita2004}). Cosmological simulations predict that most of these elusive baryons reside in a diffuse, warm-hot intergalactic medium (WHIM, hereafter) with temperatures in the range $10^5\\,\\K <T<10^7\\,\\K$ (\\citealt{Cen1999}; \\citealt{dave2001}; \\citealt{Cen2006}; see \\citealt{bertone2008} for a recent review). Theoretical models agree with numerical simulations that the relatively high temperature of the WHIM is produced by gravitational heating (\\citealt{sunyaev1972}; \\citealt{Nath2001}; \\citealt{furlanettoloeb2004}; \\citealt{rasera2006}). When large-scale structures form by gravitational collapse, cosmological shock waves are driven into the intergalactic medium (IGM) and heat the gas to temperatures much greater than $10^4\\,\\K$.  The detection of the WHIM in the low-redshift Universe requires observations in the ultraviolet (UV) and in the soft \\xray\\ bands (see e.g.\\ \\citealt{bregman2007} and \\citealt{prochaska2008} for reviews). Unfortunately, UV and \\xray\\ absorption lines are more technically challenging to detect than optical lines, and require satellites outside the Earth's atmosphere. In addition, the low density of the WHIM implies that its surface brightness is extremely low, making it nearly impossible to detect in emission with current instruments such as \\chandra\\ and \\xmm\\ and the Far Ultraviolet Spectroscopic Explorer \\fuse\\  (\\citealt{pierre2000}; \\citealt{kravtsov2002}; \\citealt{Yoshikawa2003}; \\citealt{furlanetto2004}; \\citealt{Yoshikawa2004}; \\citealt{fang2005}; \\citealt{Yoshida2005}; \\citealt{Yoshikawa2006}; \\citealt{ursino2006}). However, \\citet{werner2008} claim to have detected continuum emission with \\xmm\\ from a filament extending between the two large clusters Abell 222 and Abell 223 at $5\\sigma$ significance, thanks to the orientation of the filament along the line of sight. \\citet{zappacosta2002} and \\citet{zappacosta2005} also found evidence for diffuse thermal emission from WHIM gas in two \\rosat\\ fields, corresponding to an overdensity of galaxies at $z\\sim 0.4$ \\citep{mannucci2007} and to the Scultor Wall.  Absorption lines are less biased towards high gas densities and may therefore be more sensitive to the low-density WHIM. There are a few controversial claims of detection of \\ovii\\ absorption lines in the spectrum of Mrk 421 (\\citealt{nicastro2005}; but see \\citealt{kaastra2006}; \\citealt{rasmussen2007}) and from the Sculptor Wall \\citep{buote2009}. Future satellites such as the International \\xray\\ Observatory\\footnote{http://ixo.gsfc.nasa.gov/} (\\ixo\\ hereafter, see also \\citealt{markevitch2009} and \\citealt{bregman2009}) and \\xenia\\footnote{http://sms.msfc.nasa.gov/xenia/} \\citep{hartmann2009} will be ideal to detect a large number of WHIM absorption lines in the spectra of QSO and gamma-ray bursts, respectively (e.g.\\ \\citealt{hellsten1998}; \\citealt{perna1998}; \\citealt{kravtsov2002}; \\citealt{chen2003}; \\citealt{CenFang2006}; \\citealt{Kawahara2006}; \\citealt{paerels2008}; \\citealt{branchini2009}).  Although \\xray\\ absorption lines are a powerful tool to investigate the physical properties of the WHIM, the information they yield has the limitation of being 1-dimensional, unless close quasar pairs can be found in large numbers.  While absorption lines are not ideal to reconstruct the 3-dimensional distribution of the gas, emission lines do just that: they effectively map the 2-dimensional distribution of the WHIM on the sky, and by means of their spectral redshift they trace the distribution of the gas along the 3rd dimension. Detecting emission from the WHIM is likely to be much more challenging than detecting absorption, but the reward is comparatively large. It is therefore paramount to predict accurately what level of emission can be expected, in order to build instruments and satellites adequate to the task.  The spatial distribution of metals in the IGM, as traced by the WHIM emission, also provides important clues for understanding the metal enrichment history of the Universe. The distance travelled by metals between the sites of star formation and the lower density IGM can constrain the energy budget of processes such as galactic winds and AGN feedback.   In this work, we use a subset of runs taken from the Over-Whelmingly Large Simulations (\\owls, hereafter) project \\citep{schaye2010} to investigate which emission lines in the soft \\xray\\ band best trace the low-redshift WHIM emission. We will do the same for UV emission in a companion paper (Bertone et al., in preparation). We estimate the level of the emission in a simulated cubic region of 100 \\hm\\ comoving Mpc on a side and discuss the observability of the WHIM by future telescopes such as \\ixo, the Explorer of Diffuse emission and Gamma-ray burst Explosions (\\edge, \\citealt{piro2009}) and \\xenia.  The \\owls\\ runs \\citep{schaye2010} are well-suited for this study and represent a significant step forward with respect to previous studies. \\owls\\ couples a large simulated dynamical range with a large set of models in which the implementation of many uncertain physical processes have been varied. This special feature allows us to investigate the dependence of the results on different physical processes.  In this paper we will focus on the strongest soft \\xray\\ emission lines we have identified, which are listed in Table \\ref{eltable}, the most important being \\oviii\\ $\\lambda$ 18.97 \\AA\\ and \\ovii\\ $\\lambda$ 21.60 \\AA. When more than one line corresponds to a given transition, we consider only the strongest line.  \\begin{table} \\centering \\caption{List of the emission lines considered in this work. The first column shows the name of the emission line, the second and the third columns the wavelength and the energy of the transition, respectively. The last column identifies the type of ion responsible for the transition (e.g.\\ ``H'' means H-like) and the letter in parentheses indicates the transition for the case of triplets from He-like atoms: resonance (r), forbidden (f) and intercombination (i).} \\begin{tabular}{l r@{}l r@{}l l} \\hline \\hline Ion & $\\lambda$ \\; & (\\AA) & E \\; & (eV) & Description \\\\ \\hline \\cv & 40. & 27 & 307.  & 88 & He (r) \\\\ \\cvi & 33. & 74 & 367.  & 47 & H \\\\ \\nvi & 29. & 53 & 419.  & 86 & He (r) \\\\ \\nvii & 24. & 78 & 500.  & 24 & H \\\\ \\ovii & 21. & 60 & 573.  & 95 & He (r) \\\\ \\ovii & 21. & 81 & 568.  & 74 & He (i) \\\\ \\ovii & 22. & 10 & 560. & 98 & He (f) \\\\ \\oviii & 18. & 97 & 653.  & 55 & H \\\\ \\neix & 13. & 45 & 921.  & 95 & He (r) \\\\ \\nex & 12. & 14 & 1022. & 0 & H \\\\ \\mgxii & 8.  & 422 & 1471. & 8 & H \\\\ \\sixiii & 6.  & 648 & 1864. & 9 & He (r) \\\\ \\sxv & 5.  & 039 & 2460. & 5 & He (r) \\\\ \\fexvii & 17. & 05 & 726.  & 97 & Ne \\\\ \\hline \\hline \\end{tabular} \\label{eltable} \\end{table}  This paper is organised as follows. We briefly describe the \\owls\\ reference model and the method to calculate the line emission of the WHIM gas in Section~\\ref{owls}. Section~\\ref{emission} contains predictions for various emission lines for our reference model. Sections~\\ref{angle} and \\ref{redshift} are devoted to understanding the effects of varying the angular resolution of the detector and the redshift dependence of the flux intensity, respectively. Section~\\ref{em_from} investigates what kind of gas produces most of the emission and Section \\ref{em_model} studies the dependence of the results on the physical model. We discuss the prospects for observing soft \\xray\\ emission lines with future space-borne telescopes in Section~\\ref{detect} and we summarize our conclusions in Section~\\ref{summary}. Finally, we demonstrate in the Appendix that our conclusions are robust with respect to changes in both the numerical resolution and the size of the simulation volume. We also show in the Appendix how the results vary with the band width.  \\begin{table} \\centering \\caption{Adopted solar abundances, from \\citet{allende2001}, \\citet{allende2002} and \\citet{holweger2001}.} \\begin{tabular}{l c | l c} \\hline \\hline Element & $n_i / n_{\\rm H}$ & Element & $n_i / n_{\\rm H}$\\\\ \\hline H & 1 & Mg & 3.47$\\times 10^{-5}$ \\\\ He & 0.1 & Si & 3.47$\\times 10^{-5}$ \\\\ C & 2.46$\\times 10^{-4}$ & S & 1.86$\\times 10^{-5}$ \\\\ N & 8.51$\\times 10^{-5}$ & Ca & 2.29$\\times 10^{-6}$ \\\\ O & 4.90$\\times 10^{-4}$ & Fe & 2.82$\\times 10^{-4}$ \\\\ Ne & 1.00$\\times 10^{-4}$ & & \\\\ \\hline \\hline \\end{tabular} \\label{table_abund} \\end{table}  Unless stated otherwise, we will assume the Wilkinson Microwave Anisotropy Probe year 3 (WMAP3) $\\Lambda$CDM cosmology \\citep{spergel2007} with parameters $\\Omega_{\\rm m}=0.238$, $\\Omega_{\\rm b}=0.0418$, $\\Omega_\\Lambda=0.762$, $n=0.951$, and $\\sigma_8=0.74$. The Hubble constant is parameterised as $H_{\\rm 0} = 100~h^{-1}\\,\\kms$ Mpc$^{-1}$, with $h=0.73$. The adopted solar abundance is $Z_{\\sun} =0.0127$, corresponding to the value obtained using the default abundance set of \\cloudy\\ (version c07.02.02, last described in \\citealt{ferland1998}). This abundance set (see Table \\ref{table_abund}) combines the abundances of \\citet{allende2001}, \\citet{allende2002} and \\citet{holweger2001} and assumes $n_{\\rm He} / n_{\\rm H} = 0.1$. We note here that the abundances adopted by \\cloudy\\ may differ by significant factors from those estimated by \\citet{lodders2003}. In particular, the abundance of iron given by \\citet{lodders2003} is almost a factor of ten smaller than in \\cloudy, that of oxygen about 20 per cent higher. This should be kept in mind when comparing predictions from different studies. We stress, however, that the assumed solar abundances play no role when we use the abundances predicted by our simulations. ",
        "conclusions": "\\label{summary}  Gravitational accretion shocks  and, to a lesser extent, galactic winds driven by supernovae and/or AGN, can shock-heat the diffuse intergalactic medium to temperatures $\\ga 10^5\\,\\K$. At low redshifts a large fraction of the baryons are thought to reside in this warm-hot intergalactic medium (WHIM; $10^5\\,\\K<T<10^7\\,\\K$) and soft X-ray emission lines from metals offer a way to detect them.  Here, we used a subset of cosmological simulations from the OverWhelmingly Large Simulations (\\owls) project \\citep{schaye2010} to study the nature and detectability of this emission. These simulations are among the largest dissipative runs ever performed, and use new prescriptions for a range of baryonic processes, such as star formation, galactic winds driven by massive stars and supernovae, and feedback from AGN. Of particular relevance for us, is that they are fully chemodynamical, tracking the delayed release of 11 different elements by massive stars, supernovae of types II and Ia, and asymptotic giant branch stars, and that radiative cooling is implemented element-by-element in the presence of an evolving UV/X-ray background.  Our main conclusions can be summarised as follows: \\begin{enumerate} \\item The strongest emission in the soft \\xray\\ band comes from the \\oviii\\ line, whose emissivity peaks at a temperature of a few times $10^6\\,\\K$. A large number of other lines are, however, only slightly weaker.  \\item Emission lines from metals provide strong constraints on the temperature of the emitting gas. Most of the emitted radiation and essentially all of the potentially detectable flux comes from gas with temperatures close to the peak of the emissivity curve for a collisionally ionised gas. Because of the quadratic dependence of the emissivity on density, the emission traces neither the bulk of the baryonic nor the bulk of the metal mass. Instead, the emission is dominated by metal-rich gas at very high overdensities in and around galaxies. Since the peak temperature of the emissivity increases with atomic number and ionisation state, lines from heavier or more highly ionised atoms trace hotter gas. As such, \\cvi\\ and \\ovii\\ trace gas with $T\\sim 10^6\\,\\K$, \\oviii, \\neix\\ and \\fexvii\\ gas with $10^6\\,\\K<T<10^7\\,\\K$, and \\nex, \\mgxii, \\sixiii\\ and \\sxv\\ gas with $T\\sim 10^7$ K.  \\item By comparing results from different \\owls\\ models, we have investigated the dependence of the results on a number of physical mechanisms. It is essential to accurately model the cooling rates. In particular, studies that neglected metal-line cooling are not self-consistent and will have overestimated the maximum fluxes by about an order of magnitude. Without galactic winds, the PDF of the flux shifts to much lower values. However, the high-flux tail, which will become detectable first, is surprisingly robust to strong variations in the strength of the winds, remaining similar even when galactic winds are turned off or when violent AGN feedback is included. While assuming a constant metallicity of $0.1 Z_\\odot$ reduces the width of the flux PDF (and thus suppresses the maximum fluxes), the relative differences between the models remain similar. This suggests that the different feedback processes mainly affect the metal-line emission by changing the gas densities in and around galaxies, as the highest fluxes always come from regions with temperatures close to the peak of the emissivity curve and with (local) metallicities $\\ga Z_\\odot$.  \\item Soft \\xray\\ emission lines in the low-redshift Universe are likely to be detectable only for relatively high-density regions, such as groups and the outskirts of clusters. Because dense, compact structures can easily dominate the emission even if they contain only a very small amount of mass, it is important for future observatories to have high angular resolution. Without sufficient angular resolution, the emission is harder to detect and one risks misinterpreting the origin of the detected emission. For example, emission that looks diffuse may in reality come from a number of compact gas clouds in and around galaxies that trace the large-scale structure.  \\item While soft X-ray line emission is too faint for it to be a promising route to close the baryon budget, it does offer the exciting possibility to image the gas accreting onto and flowing out of galaxies. Currently proposed missions could provide vital clues on some of the most poorly understood, but crucial ingredients of models of the formation and evolution of galaxies. \\end{enumerate}"
    },
    "0910/0910.0106_arXiv.txt": {
        "abstract": "The neutron capture cross section of $^{14}$C is of relevance for several nucleosynthesis scenarios such as inhomogeneous Big Bang models, neutron induced CNO cycles, and neutrino driven wind models for the $r$ process. The $^{14}$C$(n,\\gamma)$ reaction is also important for the validation of the Coulomb dissociation method, where the $(n,\\gamma)$ cross section can be indirectly obtained via the time-reversed process. So far, the example of $^{14}$C is the only case with neutrons where both, direct measurement and indirect Coulomb dissociation, have been applied. Unfortunately, the interpretation is obscured by discrepancies between several experiments and theory. Therefore, we report on new direct measurements of the $^{14}$C$(n,\\gamma)$ reaction with neutron energies ranging from 20 to 800~keV. ",
        "introduction": "Inhomogeneous big bang models \\cite{ApH85} offer the possibility to bridge the mass gaps at $A=5$ and 8 and to contribute substantially to the synthesis of heavier nuclei. The suggested reaction sequence \\cite{AHS88,MaF88} for this outbreak is $^7$Li($n, \\gamma$)$^8$Li($\\alpha, n$)$^{11}$B($n, \\gamma$)$^{12}$B($\\beta^-$)$^{12}$C. Subsequent neutron captures on $^{12}$C and $^{13}$C will then lead to the production of $^{14}$C, which has a half-life of 5700$\\pm$30~yr \\cite{AjS86a}. On the time scale of big bang nucleosynthesis $^{14}$C can be considered as stable and further proton, alpha, deuteron, and neutron capture reactions on $^{14}$C will result in the production of heavier nuclei with $A \\geq 20$ \\cite{AHS88}. Due to the high neutron abundance the $^{14}$C$(n,\\gamma)$$^{15}$C reaction is expected to compete strongly with other reaction channels.  The $^{14}$C($n, \\gamma$)$^{15}$C reaction plays an important role in the discussion of neutron induced CNO cycles \\cite{WGS99} during $s$-process nucleosynthesis. Such $s$-process scenarios are characterized by comparably low neutron densities, resulting in neutron capture times, which are slow compared to typical $\\beta$ decay half lives and are associated with the He and C burning phases of stellar evolution where neutrons are produced by ($\\alpha, n$) reactions on $^{13}$C and $^{22}$Ne. While the $s$ process starts by neutron captures on iron seed nuclei, neutron captures on the light isotopes present in the burning zones can initiate a neutron induced CNO cycle. Starting from the abundant $^{12}$C, the cycle is represented by the neutron capture series on $^{12}$C, $^{13}$C, and $^{14}$C followed by the sequence $^{15}$C($\\beta^-$)$^{15}$N($n, \\gamma$)$^{16}$N($\\beta^-$)$^{16}$O($n, \\gamma$)$^{17}$O($n, \\alpha$)$^{14}$C or by producing $^{16}$N via $^{14}$N($n, p$)$^{14}$C($n, \\gamma$)$^{15}$C($\\beta^-$)$^{15}$N($n, \\gamma$). The slowest reaction in this cycle is $^{14}$C($n, \\gamma$)$^{15}$C and, therefore, $^{14}$C can build up a correspondingly high abundance. Although the outbreak from the neutron induced CNO cycle via $^{17}$O($n, \\gamma$)$^{18}$O is strongly suppressed by the dominance of the ($n, \\alpha$) channel, there might be a non-negligible effect on the neutron balance of the $s$ process depending on the cross section for $^{14}$C($n, \\gamma$). To settle this issue, the cross section would be needed for thermal energies in the 10 to 100 keV region.  Neutron capture reactions on very light nuclei carry a substantial part of the reaction flow in neutrino driven wind scenarios for the $r$ process \\cite{TSK01}. Among these reactions $^{14}$C($n, \\gamma$)$^{15}$C contributes mostly during the early phases, when $^{14}$C is formed via the sequences $^9$Be($\\alpha, n$)$^{12}$C($n, \\gamma$)$^{13}$C($n, \\gamma$)$^{14}$C and $^9$Be($n, \\gamma$)$^{10}$Be($\\alpha, \\gamma$)$^{14}$C. In these applications the ($n, \\gamma$) cross section is required up to MeV energies because of the high temperatures in excess of $3\\times 10^9$ K, corresponding to thermal energies of about 300 keV.  The $^{14}$C($n, \\gamma$) reaction is also important to validate the ($n, \\gamma$) cross sections obtained by theoretical calculations \\cite{WGT90,TBD06} via the Coulomb dissociation method in the experiments reported in Refs. \\cite{HWG02,DAB03,NFA03}. In this approach the time-reversed process is measured via breakup of $^{15}$C projectiles in the virtual photon field of a $^{208}$Pb target. The ($n, \\gamma$) cross section can then be inferred via detailed balance. The $^{14}$C($n, \\gamma$)$^{15}$C reaction is the only case with neutrons so far where both, direct and indirect, approaches were investigated experimentally. In fact, $^{14}$C belongs to the few cases where the Coulomb dissociation method can be validated in a convincingly clean way.  The first direct measurement \\cite{BWK92a} was carried out at Forschungszentrum Karlsruhe using the same sample as in the present experiment. At that time, however, the measurement was severely hampered by the fact that the nickel container used for the $^{14}$C powder sample had been strongly activated by a previous irradiation with 800~MeV protons. The present study is a repetition of this first measurement after a 12-year cooling time of the nickel container, which led to a reduction of this disturbing activity to an acceptable level. Another reason for repeating the measurement was that meanwhile a more efficient detector system for the induced activity became available. Finally, the energy range was significantly extended compared to the previous experiment.  Preliminary results of this second experiment \\cite{RHP05} turned out to be subject to significant corrections resulting from the accidental activation of the HPGe detector used. In this article we present a detailed description of the measurement as well as a thorough re-analysis of the data and of the remaining uncertainties. The results are compared to model predictions and can be used to test the applicability of the Coulomb dissociation method. ",
        "conclusions": "Compared to the result of the previous activation with $kT~=23.3~$keV \\cite{BWK92a} ($1.72~\\pm~0.43~\\mu$b) we find agreement, if the sample mass measured in this work and the currently available decay properties of \\nuc{15}{C} are taken into account. The agreement is then within 1$\\sigma$.  All available differential data for the total capture cross section of $^{14}$C are compared in Fig~\\ref{fig_results_differential}. The data are divided by $\\sqrt{E}$ to remove the energy dependence caused by the $p$-wave orbital-momentum barrier. The present cross section results are in good agreement with theoretical estimates of Wiescher {\\it et al.} \\cite{WGT90} and with the recently published estimates of Timofeyuk {\\it et al.} \\cite{TBD06} based on mirror symmetry considerations. Our data fall approximately 20\\% below the values of Descouvemont \\cite{Des00}, but exhibit the same energy dependence.  The results of Horv{\\'a}th {\\it et al.} \\cite{HWG02}, which were obtained in a Coulomb-breakup study, show a large, constant offset (Fig~\\ref{fig_results_differential}). In other words, not only the cross section values are different, but also the energy dependence. The difference can be expressed as: \\begin{eqnarray*} \\sigma_{present} &=& \\sigma_{Horvath} + c \\cdot \\sqrt{E_{c.m.}} \\\\ \\mbox{with}~c &=&0.48~\\mu\\mbox{b/keV}^{1/2} \\\\ \\end{eqnarray*} With respect to the importance of the $^{14}$C($n,\\gamma$)$^{15}$C cross section for validating the Coulomb-break-up approach for deducing this cross section from the time-reversed dissociation of $^{15}$C it is important, however, to emphasize that the present results are in good agreement with preliminary data from two other Coulomb break-up studies \\cite{DAB03,NFA03,Nak04}.   \\begin{figure} \\begin{center} \\includegraphics[width=20pc]{14c_ng_sqrt_e.eps} \\end{center} \\caption{(Color online) Comparison between the present results and previous data. \\label{fig_results_differential}} \\end{figure}  Since the paper by Beer {\\it et al.} \\cite{BWK92a}, a comparison of the differential cross section at 23.3~keV is published in most papers dealing with the \\nuc{14}C($n,\\gamma$) cross section. We note that the value published by Beer {\\it et al.} was a Maxwellian averaged cross section for $kT~=~23.3~keV$, which is different from the differential cross section at $E_{c.m.}~=~23.3~$keV. In  this tradition, a comparison of the differential 23.3~keV cross sections is presented in Fig.~\\ref{fig_results_23kev}. The present value of 5.2~$\\pm$~0.3~$\\mu$barn is based on the theoretical description of the cross section provided in the previous section.   \\begin{figure} \\begin{center} \\includegraphics[width=20pc]{14c_ng_23kev_years.eps} \\end{center} \\caption{(Color online) Comparison between this measurement (shaded band) and previous cross section results at $E_{c.m.}~=~23.3~$keV. Open squares refer to theoretical estimates while full circles refer to experiments including Coulomb-breakup studies. The only open circle refers to the measurement by Beer {\\it et al.} before the renormalization based on the new mass and line intensity information (see text). The respective references from left to right are \\cite{WGT90},\\cite{Des00},\\cite{TBD06} (theoretical) and \\cite{BWK92a}, \\cite{HWG02},\\cite{Pra02,DAB03},\\cite{NFA03,Nak04} (experimental). \\label{fig_results_23kev}} \\end{figure}  The rate suggested by \\cite{WGT90} has been used for most of the nucleosynthesis simulations of the scenarios summarized at the beginning of this paper. The agreement with the present experimental results confirms many of the previous model predictions. While present Cosmologies dismiss the likelihood of inhomogeneous Big Bang scenarios, previous simulations of the associated nucleosynthesis \\cite{RAC94} based on this $^{14}$C($n, \\gamma$) reaction rate demonstrated a substantial production of $^{14}$C at such conditions.  The role of the $^{14}$C($n, \\gamma$)$^{15}$C reaction as the slowest link in the neutron induced CNO cycles proposed by \\cite{WGS99} is also confirmed by the present results. Detailed simulations now help to analyze the impact of such a cycle on the neutron flux during core carbon burning and shell carbon burning. These results indicate that many more branches exist due to the presence of charged particles in stellar helium and carbon burning environments \\cite{PGH07}. For helium burning most of the $^{13}$C produced by $^{12}$C($n, \\gamma$) is depleted by the $^{13}$C($\\alpha, n$) reaction rather than by $^{13}$C($n, \\gamma$) and the production of $^{14}$C is negligible as shown already by \\cite{TAG96}. This may be different for shell carbon burning which is characterized by higher $^{12}$C abundances and a significantly lower $\\alpha$ flux. New simulations on aspects of neutron production and capture reactions are presently in preparation \\cite{PGW08}. The study indicates that the main production of $^{14}$C is given by the two reactions $^{14}$N($n, p$)$^{14}$C and $^{17}$O($n, \\alpha$)$^{14}$C. Because of the here confirmed low cross section, the $^{14}$C($n,\\gamma$) reaction does not play a significant role for reducing the $^{14}$C abundance. However, because of the relatively high temperatures of T$\\approx$1 GK in the carbon burning zone, alternative depletion channels open via $^{14}$C($p, n$)$^{14}$N with a negative Q-value of -626 keV and via $^{14}$C($\\alpha,\\gamma$)$^{18}$O alpha capture providing a new abundance balance.  New simulations are also underway for studying the impact of neutron capture reactions on neutron rich Be, B, and C isotopes on the nucleosynthesis of light elements in neutrino driven wind supernova shock scenarios \\cite{BGM06}. The completion of these studies does however require a detailed analysis of neutron capture reactions on short-lived neutron rich isotopes to simulate the anticipated reaction flow reliably \\cite{TSK01}. New shell model based simulations of these rates are presently in preparation taking also into account the rapidly growing experimental nuclear structure information on neutron rich nuclei in the Be to Ne range."
    },
    "0910/0910.4725_arXiv.txt": {
        "abstract": "Models for the steady state collisional evolution of low eccentricity planetesimal belts identify debris disks with hot dust at 1AU, like $\\eta$ Corvi and HD69830, as anomalous since collisional processing should have removed most of the planetesimal mass over their $>1$ Gyr lifetimes. This paper looks at the effect of large planetesimal eccentricities ($e \\gg 0.3$) on their collisional lifetime and the amount of mass that can remain at late times $M_{\\rm{late}}$. Assuming an axisymmetric planetesimal disk with common pericentre distances and eccentricities $e$, we find that $M_{\\rm{late}} \\propto e^{-5/3}(1+e)^{4/3}(1-e)^{-3}$. For a scattered disk-like population (i.e., with common pericentre distances but range of eccentricities), in the absence of dynamical evolution, the mass evolution at late times would be as if only planetesimals with the largest eccentricity were present in the disk. Despite the increased remaining mass, higher eccentricities do not increase the amount of hot emission from the collisional cascade until $e>0.99$, partly because most collisions occur near pericentre thus increasing the dust blow-out diameter. However, at high eccentricities ($e>0.97$) the blow-out population extending outwards from pericentre may be detectable above the collisional cascade; higher eccentricities also increase the probability of witnessing a recent collision. All of the imaging and spectroscopic constraints for $\\eta$ Corvi can be explained with a single planetesimal population with pericentre at 0.75AU, apocentre at 150AU, and mass $5M_\\oplus$; however, the origin of such a high eccentricity population remains challenging. The mid-infrared excess to HD69830 can be explained by the ongoing destruction of a debris belt produced in a recent collision in an eccentric planetesimal belt, but the lack of far-infrared emission would require small bound grains to be absent from the parent planetesimal belt, possibly due to sublimation. The model presented here is applicable wherever non-negligible planetesimal eccentricities are implicated and can be readily incorporated into N-body simulations. ",
        "introduction": "\\label{s:intro} A natural byproduct of the planet formation process, at least in the core accretion model, is the formation of planetesimals (Lissauer 1993). Evidence for planetesimals following the protoplanetary disk phase comes from debris disks, a phenomenon in which main sequence stars exhibit an infrared excess which is attributed to the thermal emission of dust released from planetesimals in collisions and sublimation (see review in Wyatt 2008). The Solar System has its own debris disk, the majority of which is in the asteroid and Kuiper belts.  Typically extrasolar debris disks are observed to lie in a ring at a single radius (Greaves et al. 2005; Kalas, Graham \\& Clampin 2005; Schneider et al. 2009), or where they are not imaged the emission spectrum is dominated by a single temperature (Chen et al. 2006). This motivates considering these disks as planetesimal belts that are directly analogous to the asteroid and Kuiper belts (Moro-Mart\\'{i}n et al. 2008), and the disks where dust is detected at multiple radii (Wyatt et al. 2005; Absil et al. 2006; Smith et al. 2009a; Backman et al. 2009; Chen et al. 2009) are usually inferred to have multiple planetesimal belts. In the absence of other dynamical processes, the evolution of these belts is expected to be dominated by collisions which grind away the mass of the largest objects into dust which is subsequently removed by radiation pressure (or P-R drag in the case of the Solar System) (Wyatt 2009).  The collisional evolution of the planetesimal belts of the Solar System has been studied extensively. Collision rates can be derived accurately between objects moving on given orbits (\\\"{O}pik 1951; Wetherill 1967; Greenberg 1982; Bottke et al. 1994; Vedder 1996; Dell'Oro \\& Paolicchi 1998), and the steady state size distribution of the belts resulting from their collisional evolution is both well understood analytically (Dohnanyi 1969; Tanaka et al. 1996; O'Brien \\& Greenberg 2003; Kobayashi \\& Tanaka 2009) and one-dimensional numerical models of this evolution that include a realistic prescription for the outcome of collisions provide a good fit to the observed size distributions in the asteroid belt (Davis et al. 1989; Durda, Greenberg \\& Jedicke 1998; Bottke et al. 2005; O'Brien \\& Greenberg 2005) and Kuiper belt (Stern \\& Colwell 1997; Davis \\& Farinella 1997; Kenyon \\& Bromley 2004).  The approach to considering the collisional evolution of extrasolar debris disks is slightly different in that the orbital element and size distributions of the parent bodies are poorly constrained, rather it is important to generalize the effect of this evolution on debris disk observability with respect to parameters such as initial planetesimal belt mass, radius and mean eccentricity. Analytical models that achieve this were developed by Dominik \\& Decin (2003) who considered the evolution of a mono-disperse planetesimal belt (i.e., with planetesimals all of the same size) that feeds a population of smaller planetesimals and dust that is observed. This model was later refined by Wyatt et al. (2007a) to consider the parent planetesimals and the smaller objects to form a continuous size distribution defined by a single power law as expected for the steady state case where planetesimal strength is independent of size (Dohnanyi 1969; Tanaka et al. 1996). A size dependent planetesimal strength was later included in such models by L\\\"{o}hne et al. (2008) resulting in a more realistic 3 phase size distribution (e.g., O'Brien \\& Greenberg 2003). Both the Wyatt et al. (2007b) and L\\\"{o}hne et al. (2008) models were applied to the statistics of detections of debris disks around A stars and Sun-like stars to show that these could be explained if the majority of such debris disks evolve purely due to steady state collisional evolution.  One important result that came out of the Wyatt et al. (2007a) study was the concept of a maximum planetesimal belt mass, and hence a maximum dust luminosity, that can remain for a given radius belt at a given time, regardless of its initial mass. Although this is no longer strictly true when a size dependent strength is used, L\\\"{o}hne et al. (2008) showed that initial mass has a relatively modest effect on the mass remaining at late times, and concurred that for realistic planetesimal belt parameters there is indeed a maximum planetesimal belt mass and dust luminosity for a given age (see also Heng \\& Tremaine 2009). This concept was used by Wyatt et al. (2007a) to show that 1-2Gyr systems like $\\eta$ Corvi and HD69830 that have large quantities of hot dust at $\\sim 1$AU (Wyatt et al. 2005; Beichman et al. 2005; Smith, Wyatt \\& Dent 2008), cannot be replenishing that dust from planetesimal belts that are coincident with the dust (i.e., analogous asteroid belts). They concluded that the parent bodies of the observed dust must have originated at larger radii ($\\gg$ several AU) where collisional processing times would have been longer. The paper also concluded that the hot dust is transient and proposed that this might have been scattered in from an outer belt in an epoch akin to the Late Heavy Bombardment in the Solar System (see review in Hartmann et al. 2000). There are now several examples of systems exhibiting hot dust that appears to be transient by the criterion described by Wyatt et al. (2007a) (e.g., di Folco et al. 2007; Akeson et al. 2009; Mo\\'{o}r et al. 2009).  The motivation of this paper is to consider whether it is possible to circumvent the conclusion that the hot dust in systems like $\\eta$ Corvi and HD69830 must be transient by postulating a population of parent planetesimals on highly eccentric orbits ($e \\gg 0.3$). In such a model the hot dust would originate from material close to pericentre, and the parent population could be long-lived because the planetesimals spend most of their time at apocentre. This would challenge our traditional view of debris disks as belts of planetesimals with modest eccentricity ($e<0.3$), which is also implicit in the models of Wyatt et al. (2007a) where collision velocities are assumed to be proportional to Keplerian velocity times a mean eccentricity for the belt, and in the models of L\\\"{o}hne et al. (2008) where eccentricities up to 0.35 were considered. However, it is clear from the Solar System that there are also populations of planetesimals on highly eccentric orbits ($e \\gg 0.3$): the comets scattered in from the Kuiper belt (Duncan 2008); the scattered disk of the Kuiper belt (which may be primordial in origin and extends all the way to the Oort cloud) (Gomes et al. 2008); as well as the Near Earth Asteroids (Bottke et al. 2002). While the contribution of these populations to the dust content of the zodiacal cloud may be small, the cometary contribution could be as much as 90\\% (Ipatov et al. 2008; Nesvorn\\'{y} et al. 2009), and may have been significantly higher in the past, e.g., during the epoch known as the Late Heavy Bombardment (Gomes et al. 2005; Booth et al. 2009). Furthermore the opposite may be true for planetary systems with different architectures and formation scenarios, in which eccentric planetesimals may dominate. Indeed, planet formation models often predict a highly eccentric remnant planetesimal population (Edgar \\& Artymowicz 2004; Mandell, Raymond \\& Sigurdsson 2007; Payne et al. 2009).  Thus, here we develop the model of Wyatt et al. (2007a) to include interactions between planetesimals of arbitrary eccentricities and semimajor axes (and inclinations). Although this model does not (yet) include the more realistic assumption of a size dependent planetesimal strength, it benefits by providing simple analytical formulae for collision lifetimes from which the observability of a planetesimal belt as a function of its eccentricity can be readily assessed. The inclusion of a size dependent strength would be expected to give results within an order of magnitude of those presented here (see, e.g., Fig. 11 of L\\\"{o}hne et al. 2008), a level of uncertainty that is commensurate with the uncertainty in estimates for planetesimal strength at each size for the Solar System and for different assumptions about planetesimal composition (see e.g., Fig. 1 of Durda et al. 1998, and Fig. 11 of Leinhardt \\& Stewart 2009).  In \\S \\ref{s:2} we consider the collisional evolution of an axisymmetric disk of planetesimals all of which have the same pericentre and apocentre distances, and show that the concept of a maximum remaining mass for a given age also applies in this case, but that the remaining mass is higher for higher eccentricities (if pericentre distance is kept constant). The approach to calculating collision rates is similar to that of Bottke et al. (1994) in that we assume random mean longitudes, arguments of pericentre and longitudes of ascending node, but differs in using a particle-in-a-box approach to calculate the collision rate at a particular point on the orbit then integrating around the orbit (as opposed to calculating this from the fraction of the orbits that the planetesimals spend close enough that they overlap in physical space). Our collision rate at each point also includes an integration over the size distribution of impactors that can cause a catastrophic collision, whereas this integration is performed after calculating the intrinsic collision probability by Bottke et al. (1994), requiring that method to keep track of the velocity probability distribution. Our integration is performed using a Monte-Carlo technique, but in the case of mutual collisions amongst a population with common eccentricities and semimajor axes, and assuming that collision velocities are dominated by radial motion (due to eccentricities) rather than vertical motion (due to inclinations), the collision rate can also be derived analytically.  In \\S \\ref{s:4} we consider the more general situation in which planetesimals can interact with planetesimals with different pericentre and apocentre distances, and show that our collision rates agree with those of the most accurate studies available in the literature. To consider the evolution of a realistic planetesimal belt where a range of eccentricities and semimajor axes is present we adopt an approach similar to Krivov et al. (2005, 2006) in that we consider the evolution of the phase space distribution.\\footnote{Note that the Krivov et al. 2005 model is not accurate for high eccentricities because of the way mean impact velocity was calculated (see \\S \\ref{ss:comp}).} Here we outline a scheme for evolving the phase space distribution numerically, and apply this to a scattered-disk like population with common pericentre and a range of apocentre distances. The emission properties of eccentric rings are considered in \\S \\ref{s:3} to assess if the emission spectrum of real systems can be consistently explained by steady state processing given the stellar ages. The conclusions are given in \\S \\ref{s:conc}. ",
        "conclusions": "\\label{s:conc} This paper considers collisional processes in a population of planetesimals with high eccentricities. Collision rates are derived both analytically and using Monte-Carlo simulations. It was found that eccentricity has a significant effect on collision rates, and that the amount of mass that can remain in a planetesimal belt at late times can be significantly increased.  The emission properties of eccentric planetesimal belts were presented, and it was found that radiation pressure causes eccentric rings to be deficient in small particles. Thus, despite the increased mass of a high eccentricity planetesimal belt at late times, extreme eccentricities of $>0.99$ are required to enhance the emission level above that expected from a low eccentricity belt. However, the high mass loss rate of extreme eccentricity planetesimal belts can cause the wind of blow-out particles that extends outward from the pericentre to be detectable. The high frequency of massive collisions in such belts can also make it likely for us to be witnessing dust produced in such collisions.  Application of this model to the $\\eta$ Corvi debris disk showed that all available observations can be explained by an extreme eccentricity ($e=0.99$) planetesimal belt of mass $5M_\\oplus$, circumventing the conclusion that the hot dust at 1.7AU must be transient. Despite this success, the dynamical challenges of creating such a massive extreme eccentricity population would need to be overcome before this model can be considered viable. Observational tests are suggested including the presence of resolvable emission (and absence of planets) in the 3-100AU region. Application to HD69830 is complicated by the lack of far-infrared emission. It may be possible to explain this non-detection by the removal of small dust from the collisional cascade by sublimation, in which case the mid-infrared emission may be plausibly explained by the ongoing destruction of a debris belt produced in a recent collision in an eccentric planetesimal belt.  Although the majority of the discussion focuses on single eccentricity populations, the results can also be applied to populations with a range of semimajor axes and eccentricities. This was demonstrated by application to scattered disk-like populations where it was found that, in the absence of other dynamical processes, rapid collisional erosion of the low eccentricity populations would be expected to result in a single high eccentricity population. Since the known high eccentricity planetesimal populations are produced in interactions with planets, and so may be continuing to undergo dynamical evolution on timescales shorter than collisional timescales, it is noted that dynamical interactions may continue to play a defining role in the long term evolution of high eccentricity populations, and that the collisional evolution scheme presented here could be readily incorporated into N-body simulations of planet-planetesimal interactions to derive simultaneously the collisional and dynamical evolution of a scattered planetesimal population. A further extension of the model would include a prescription for planetesimal strength as a function of size which would lead to a departure from the single phase size distribution assumed here (e.g., L\\\"{o}hne et al. 2008).  The results of this study would be applicable wherever non-negligible planetesimal eccentricities are found. Thus, other potential applications include the Solar System's comet and NEO populations, particularly in the early phases when these populations would have been more massive and hence collisional processing more important (Booth et al. 2009), the outcome of planet formation models (e.g., Payne et al. 2009), and systems where eccentric planetesimals may be implicated such as the origin of dust around White Dwarfs (Farihi, Jura \\& Zuckerman 2009) and of the hottest dust population of debris disks like Vega (Absil et al. 2006)."
    },
    "0910/0910.3043_arXiv.txt": {
        "abstract": "Recent advances in theory and computational experiments have shown the need to refine the previous categorisation of magnetic reconnection at three-dimensional null points -- points at which the magnetic field vanishes. We propose here a division into three different types, depending on the nature of the flow near the spine and fan of the null. The spine is an isolated field line which approaches the null (or recedes from it), while the fan is a surface of field lines which recede from it (or approach it).  So-called {\\it torsional\\ spine\\ reconnection} occurs when field lines in the vicinity of the fan rotate, with current becoming concentrated along the spine, so that nearby field lines undergo rotational slippage.  In {\\it torsional\\ fan\\ reconnection} field lines near the spine rotate and create a current that is concentrated in the fan with a rotational flux mismatch and rotational slippage. In both of these regimes, the spine and fan are perpendicular and there is no flux transfer across spine or fan.  The third regime, called {\\it spine-fan reconnection}, is the most common in practice and combines elements of the previous spine and fan models. In this case, in response to a generic shearing motion, the null point collapses to form a current sheet that is focused at the null itself, in a sheet that locally spans  both the spine and fan. In this regime the spine and fan are no longer perpendicular and there is flux transfer across both of them. ",
        "introduction": "\\label{sec:1} Magnetic reconnection is a fundamental process of energy release that lies at the core of many dynamic phenomena in the solar system such as solar flares, coronal heating events, geomagnetic substorms and flux transfer events. Reconnection in three dimensions has been shown to be completely different in many fundamental respects from the classically studied process in two dimensions \\citep{schindler88,hesse88,priest03}. The main thrust of reconnection theory at present is to understand the different ways in which it may take place in three dimensions \\citep[e.g., the books][]{priest00,birn07}. A key point is that in three dimensions reconnection occurs where a component of the electric field parallel to the magnetic field is present -- and this can be in many different field configurations. For example, reconnection may occur either at null points \\citep[e.g.,][]{craig95,priest96a,craig98,craig00,hornig03,heerikhuisen04,pontin05a,pontin06a} or in the absence of null points at quasi-separatrix layers or hyperbolic flux tubes \\citep{priest95a,demoulin96a,demoulin96b,demoulin97d,hornig98,titov03a,hesse05,demoulin06a,titov07a,titov09} or it may occur along separators that join one null point to another \\citep{priest96a,longcope96a,galsgaard96b,galsgaard00b,longcope01,parnell04,priest05a,longcope05,parnell08a}.  Null points are common in the solar atmosphere \\citep{filippov99,schrijver02,longcope03,close04c} and are sometimes implicated in solar flares and coronal mass ejections \\citep{longcope96b,fletcher01,aulanier00,aulanier05a,aulanier06a,cook09}. Three-dimensional collapse of a null has been described \\citep{bulanov84,bulanov97,klapper96,parnell96,parnell97} and stationary resistive flows near them have been modelled \\citep{craig96,craig99,titov00}. In particular, for a linear null and uniform magnetic diffusivity, \\citet{titov00} discovered field-aligned flows when the spine current is small and spiral field-crossing flows which do not cross the spine or fan when the spine current exceeds a critical value.  A three-dimensional null point possesses two different classes of field lines that connect to the null: for a so-called {\\it positive null point}, a surface of field lines (called a {\\it fan} by \\citet{priest96a}) recede from the null, while an isolated field line (called the {\\it spine} of the null) approaches  it from two directions; for a {\\it negative null point}, on the other hand the fan approaches the null, while the spine recedes from it. (For an alternative nomenclature see Ref.~\\cite{lau90}.) The different types of linear null were categorised by \\citet{parnell96}. The generic null in a potential magnetic field is an improper radial null, with the fan perpendicular to the spine and the field lines in the fan approaching or receding from essentially two directions (Fig.~\\ref{fig:1}b). A particular case is the proper radial null in which the field lines in the fan are radial (Fig.~\\ref{fig:1}a). The effect of a current along the fan is to make the fan and spine no longer perpendicular (Fig.~\\ref{fig:3}b), whereas a strong enough current along the spine makes the fan field lines spiral (Fig.~\\ref{fig:3}a).  There have been three steps towards categorising reconnection at a null point due to (i) analytical ideal modelling, (ii) kinematic resistive modelling and (iii) computational experiments.  The initial analytical ideal treatment by \\citet{priest96a} aimed to understand the types of ideal motions that are possible in the environment of a null point. They supposed that the nature of reconnection is determined to a large extent by the nature of the large-scale flows: they suggested that an ideal flow across the fan would drive {\\it spine\\ reconnection}, in which a current forms along the spine, whereas an ideal flow across the spine would drive {\\it fan\\ reconnection} with a strong current in the fan. They also proposed {\\it separator\\ reconnection} with a strong current along a separator joining two nulls.  Since then, as we shall see in this paper, although behaviour reminiscent of the early spine and fan models may be observed in certain limiting situations, recent numerical experiments have suggested different forms of spine and fan reconnection and also a hybrid spine-fan regime as being the generic modes that occur in practice. However, the existence of separator reconnection has been well confirmed by a series of numerical experiments \\citep{galsgaard96b,galsgaard00a,parnell04,parnell08a} and its importance in the solar corona has been stressed \\citep{longcope01,longcope05}.  In addition, quasi-separatrix layer reconnection (called slip-running reconnection by \\citet{aulanier06a}) has been confirmed in numerical experiments \\cite{birn00,pontin05c,aulanier05a,mellor05,demoortel06a,demoortel06b,wilmot07a} and in bright point and flare simulations \\citep{demoulin93b,demoulin97a,demoulin97d,masson09a,torok09}.  Our aim here is simply to look more closely at the nature of reconnection at a 3D null point and to propose a new categorisation to replace spine reconnection and fan reconnection.  In the next section it is necessary to summarise the main results from  theory and computational experiments on null-point reconnection and to reinterpret them in the light of the new regimes of reconnection that we are proposing. In the following sections we consider in turn the properties of the three new types of reconnection, namely, {\\it torsional\\ spine\\ reconnection}, {\\it torsional\\ fan\\ reconnection} and the most common regime {\\it spine-fan\\ reconnection}. ",
        "conclusions": "\\label{sec:6} We have here outlined a new categorisation of reconnection regimes at a 3D null point.  In place of the two previous types, namely, spine and fan reconnection, we suggest that three distinct generic modes of null point reconnection are likely to occur. The first two are caused by rotational motions, either of the fan or of the spine, leading to either {\\it torsional spine reconnection} or {\\it torsional fan reconnection}.  These involve slippage of field lines in either the spine or the fan, which is quite different from classical 2D reconnection, but does involve a change of magnetic connection of plasma elements.  Even though pure spine or fan reconnection may occur in special situations (such as when $\\nabla \\cdot {\\bf v}=0$ or  there are high-order currents), it is much more likely in practice that a hybrid type of reconnection takes place that we refer to as {\\it spine-fan reconnection}. This is the most common form of reconnection that we expect to see in three dimensions at a null point, since it is a natural response to a shearing of the null point. It is most similar of all the 3D reconnection regimes to classical 2D reconnection and involves transfer of magnetic flux across both the spine and the fan. It possesses a diffusion region in the form of a current sheet that is inclined to the fan and spine and has current localised in both the spine and fan, focussed at the null.  In future, much remains to be determined about these new regimes of reconnection that have been observed in numerical experiments. One is the shape and dimensions of the diffusion regions and their relation to the driving velocity and the magnetic diffusivity.  Another key question is: what is the rate of reconnection at realistic plasma parameters, and is there a maximum value?  Since the analytical theory is so hard in three dimensions, progress is likely to be inspired by future carefully designed numerical experiments."
    },
    "0910/0910.4380_arXiv.txt": {
        "abstract": "We have studied the dark matter (DM) distribution in a  $\\approx 10^{12} h^{-1}  M_{\\sun}$ mass halo extracted from a simulation consistent with the concordance cosmology, where the physics regulating the transformation of gas into stars was allowed to change producing galaxies with different morphologies.  The presence of baryons produces the  concentration of the DM halo with respect to its corresponding dissipationless run, but we found that this response does not only depend on the amount of baryons gathered in the central region but also on the way they have been assembled. DM and baryons affect each other in a complex way so the formation history of a galaxy plays an important role on its final total mass distribution. Supernova (SN) feedback regulates the star formation and triggers galactic outflows not only in the central  galaxy but also in its satellites. Our results suggest that, as the effects of SN feedback get stronger, satellites get less massive and can even be more easily disrupted by dynamical friction, transferring less angular momentum. We found indications that this angular momentum could be acquired not only by the outer part of the DM halo but also by the inner ones and  by the stellar component in the central galaxy. The latter effect produces stellar migration which contributes to change the inner potential well, probably  working against further DM contraction.  As a consequence of the action of these processes, when the  halo hosts a galaxy with an important disc structure formed by smooth gas accretion, it is more concentrated than when it hosts a spheroidal  system which experienced more massive mergers and interactions. We also found that in the later case, the halo has less radial velocity anisotropy than when the halo hosts a disc galaxy. In most of our runs with baryons, we do not  detect the inversion of the velocity dispersion characteristic of the dissipationless haloes. We have found that rotation velocities for the systems that were able to develop a disc structure  are in good agreement with the observations and none of them have been formed satisfying the adiabatic contraction hypothesis. ",
        "introduction": "Over the last  decades numerical simulations have became a powerful tool to study the validity of cosmological models. $\\Lambda$-CDM scenarios have been found to be able to reproduce successfully the global properties of the observed structure  at large scales. However, at galactic scale, the so-called concordance ($\\Lambda$-CDM) paradigm has been challenged by several  observational results. In this respect, the cuspy inner profile obtained in many of the numerical simulations has been claimed  not to be consistent with the rotation curves observed for low surface brightness (Flores \\& Primack 1994; Moore 1994; Dutton, van den Bosch \\& Courteau 2008; Salucci, Yegorova \\& Drory 2008) and dwarf galaxies (e.g. Gnedin \\& Zhao 2002). There have been many attempts to explain the possible erasement of the inner cusp via different mechanisms: interactions through dynamical friction with the substructure (e.g. El-Zant et al. 2001, Tonini, Lapi \\& Salucci 2006), SN feedback  (Mashchenko et al. 2006) etc., but the problem is not yet solved. \u00b7  Another important  pending issue is the overabundance of small dark matter (DM) subhaloes (e.g. Moore et al. 1999; Stadel et al. 2009). The total number of subhaloes found in the simulations is much larger than the number of known satellite galaxies surrounding the Milky Way. The inner shape of the DM  density profiles and the abundance of subhaloes are of particular interest because they can be  used to perform several observational diagnostics such as gravitational lensing (Zackrisson \\& Riahm 2009). Also, the prospect of detecting DM particles annihilation makes it essential a deep understanding of the small scale distribution of the DM in the central regions of our Galaxy (Diemand et al. 2008, Springel et al. 2008).  Although the ``universality'' of the spherically-averaged density profiles (Navarro, Frenk and White 1996, hereafter NFW) of $\\Lambda$-CDM  haloes has gained a broad consensus over the past couple of decades, new evidences based on high resolution simulations have been found against it (Navarro et al. 2004, Merrit et al. 2006, Gao et al. 2008). Merrit et al. (2006) tested different fitting functions and found that the profiles were better described by a de Vaucouleurs' law instead of the NFW two parameter formula. They also found a systematic variation in the profile shape with halo mass. Navarro  et al. (2008, hereafter N08) analysed different galaxy-sized haloes simulated with unprecedent high resolution and found small but significant deviations from the so-called NFW universal profile.  The contraction of the DM haloes due to the infall and condensation of baryons in the central regions is a well accepted process (e.g. Barnes \\& White 1984). Many attempts to predict their effects have been made through  models based on the adiabatic contraction (AC) hypothesis  such as the one developed by Blumenthal et al. (1986, hereafter B86). However, it has been shown that this kind of models overestimates the level of contraction. More recently,  Gnedin et al. (2004) and Sellwood \\& McGaugh (2005) developed AC models based on the work of Young (1980) which considered the possibility of radial motions. However, all  AC-based models  missed the hierarchical characteristic of the assembly of the galaxy as suggested by the current large scale observations (e.g. S\\'anchez et al. 2006), which might have non negligible consequences on the final distribution of the DM (e.g. Debatistta et al. 2008).  In a fully cosmological context, simulations  have already shown how the DM haloes concentrate  when  baryons are included, providing hints of a possible dependence on the assembly history  (e.g. Tissera \\& Dom\\'{\\i}nguez-Tenreiro et al. 1998; Gnedin et al. 2004; O\\~norbe et al. 2007). Recently Romano-D\\'{\\i}az et al. (2008) analysed the evolution of the central DM profile   in cosmologically grown galactic haloes, claiming that  when baryons are present the cusp is gradually levelled off. They suggest that this effect could be associated with the action of the subhaloes that heat up the cusp region of the DM halo through dynamical friction, and force it to expand (Ma \\& Boylan-Kolchin 2004; Debatistta et al. 2008).  In order to help shed light on these issues, we study a set of intermediate resolution cosmological simulations where different baryonic structures have been able to  form from identical initial conditions via the modification of the physics of baryons. The set of simulations studied in our work are those analysed by Scannapieco et al. (2008, hereafter S08) where  the star formation activity and the SN feedback were modified in order to study the role played by each of these processes. As a consequence, a variety of galaxies were obtained, each one exhibiting different morphological and dynamical properties. These experiments allow us to analyse how the dark matter evolves when baryons are assembled in a different fashion but governed by the same underlying merger tree. The first results of this analysis were reported by  Pedrosa, Tissera \\& Scannapieco (2009) where it is clearly shown that the final structure of the dark matter halo  depends on the way baryons are put together and not solely on the amount of baryons gathered in the centre. We also found hints for a  re-distribution of angular momentum related to the accretion of satellites. In this paper, we extend this work and analyse in detail the DM haloes considering the evolution of baryons.    This paper is organized as follows. In Section 2, we describe the numerical experiments and summarize the main features of the simulated galaxies. In Section 3, we analyse the DM density profiles. In Section 4, we study the interaction with satellites. In Section 5, rotation curves and the AC hypothesis are analysed. In Section 6 we summarize our main results. ",
        "conclusions": "We have studied the DM distribution in a set of runs of a  $\\approx 10^{12} M_{\\sun}$ halo extracted from a cosmological simulation, where the physics that regulates the  SF activity and SN feedback was varied allowing the formation of galaxies with different morphologies at $z=0$. Since the underlying DM merger tree is the same in all our runs, the differences in the properties of the DM and baryons can be directly ascribed to the variations in the baryonic physics. For the same reasons, all simulations have been run with the same numerical resolution so they are affected by resolution in a similar way. Because we are studying the DM distribution principally in the central regions, we have estimated all quantities outside three gravitational softenings. And as a further test of the robustness of our results against numerical artifacts, we have studied the DM profiles in two haloes selected from a fully cosmological simulation with higher numerical resolution. These haloes host galaxies with different morphologies and reproduce remarkably well  our findings (appendix A).  Our main results can be summarized as follows:  \\begin{enumerate}  \\item[1. ]We found that the Einasto's model provides the best fit for the spherically-averaged density profiles of our DM haloes. When baryons are present, the haloes become more concentrated in the central regions. However, the amount of baryons collected within the inner regions does not by itself  determine the response of the DM halo to the assembly process of the galaxy.  \\item[2. ]When baryons are included, the velocity dispersion increases in the central region compared to the dissipationless case and no ``temperature inversion'' is observed, except for the  E-3 run where the fraction of baryons remaining in the halo is very small due to its strong SN feedback. The slope of the inner velocity dispersion profiles increases with increasing baryonic mass collected at the centre. We found that haloes hosting spheroidal galaxies tend to have weaker levels of velocity anisotropy than the DM-only run. While those haloes with an important disc galaxy show the highest levels of velocity anisotropy.  \\item[3. ]The formation history plays an important role in the final distribution of its DM halo. We observed that those systems that are able to develop inside-out-formed discs, although they host in the central regions a lower amount of baryons than  galaxies that formed  old extended spheroids, have  more concentrated DM  profiles. Since all our runs shared the same merger tree, the differences between them can be directly ascribed to their different baryonic evolutions  which are determined mainly by the SN feedback.  \\item[4. ]We followed the evolution of the DM distribution in the central regions with redshift via the concentration parameter $\\Delta_{v/2}$. We found that all haloes increase their concentration as they grow in time. As expected, the dissipationless run has the lowest concentration at all times. Also, we observed that  haloes present a flattening in this relation with respect to the critical relation indicating an expansion of the mass distribution in the central regions. This flattening differs among the different haloes being larger for those systems hosting an spheroidal galaxy. This trend can be linked with close approaches of satellites and their properties.  \\item[5. ] The analysis of  the satellites within the virial radius of the progenitor objects as a function of redshift yields that in the NF case, the satellites are clearly more massive and have been able to survive further in the halo, since stars are more gravitationally bounded. The systems that, at lower redshifts, developed a disc component (e.g E-0.7 and F-0.9) show less massive satellites at all redshifts as  a consequence of the action of SN feedback. As expected, the most diffuse satellites correspond to the E-3 case. By analysing the mass of the satellites within the virial radius, we found noticeable mergers or disintegration episodes which can be correlated with features in the specific angular momentum content of the mass components.  We studied the specific angular momentum content of main baryonic component and its halo from  $z \\approx 1.6$, when the $\\Delta_{v/2}$ starts to clearly indicate an expansion of the central DM concentration. We selected the  NF and E-0.7 haloes as case studies. For these cases, we found that even the inner $10\\%$ of the stellar mass gains angular momentum as a function of redshift, although the increase is more important for the NF case. The gas component shows a similar behaviour but, in this case, it is in the E-0.7 run where the gas acquired larger amount of angular momentum  induced by the SN feedback which we know is successful at driving galactic outflows (S08). The DM halo (without substructure) increases its angular momentum content at all mass bins. Again, the larger profits are measured for the NF case which has the most massive orbiting satellites. Hence, the mass components identified at $ z \\approx 1.6$  in NF case  gained more angular momentum up to $z=0$ than their counterparts in E-0.7, except for the gas component which is anyway less massive than the other ones. The fact that stars migrate also contributes to change the inner potential well since they are the dominating mass component in the central region. This, on its turn,  probably acts against further DM contraction.  Both effects could explain  the evolution of $\\Delta_{v/2}$ and the fact that, when the halo hosts a disc-dominated galaxy, it is more concentrated than when it hosts a spheroid-dominated one.   \\item[6. ]The baryonic rotation curves of our simulated galaxies reflect the presence of a concentrated dominating spheroid or a dominating disc component. The total circular velocity in the central regions is mainly determined by the baryonic component. When an important disc component is present the total velocity distribution gets flat in the central regions. The evolution of the baryonic circular velocity with time of the E-0.7 run shows that as the disc forms, the curve became flatter out to larger radii. We  quantified the flattening of the curve through the LS and correlated it with the shape parameter. Galaxies with lower values of LS tend to have haloes with larger $n$ parameter. We found values for $\\frac{V_{max}}{V_{200}} $ between $\\approx 1.15$ and $1.5$. The systems where a disc galaxy was able to form present the lowest ratios, in agreement with observational results.  \\item[7. ]We have compared our simulated haloes with different AC prescriptions. We found, as expected, that the Blumenthal et al (1986) model overpredicts the level of concentration and also changes the shape of the DM distribution. The recipes of Gnedin et al. (2004) and A09 are an  improvement over the B86 approach. However, they overpredict the level of contraction when the haloes host an spheroid-dominated galaxy. From our analysis, we have obtained a prescription that provides a better representation for the contraction of our haloes, depending on the morphology of the galaxy. However, the 'universality' of this prescription should be tested with a larger statistical sample.    All our findings indicate that the response of the DM halo to the presence of baryons is the result of the joint evolution of baryons and DM during the assembly of the galaxy and in this sense,  the cosmological context for galaxy formation cannot be ignored.  \\end{enumerate}"
    },
    "0910/0910.3782_arXiv.txt": {
        "abstract": "We present results on the compact steep-spectrum quasar 3C~48 from observations with the Very Long Baseline Array (VLBA), the Multi-Element Radio Linked Interferometer Network (MERLIN) and the European VLBI Network (EVN) at multiple radio frequencies. In the 1.5-GHz VLBI images, the radio jet is characterized by a series of bright knots. The active nucleus is embedded in the southernmost VLBI component A, which is further resolved into two sub-components A1 and A2 at 4.8 and 8.3 GHz. A1 shows a flat spectrum and A2 shows a steep spectrum. The most strongly polarized VLBI components are located at component C $\\sim$0.25 arcsec north of the core, where the jet starts to bend to the northeast. The polarization angles at C show gradual changes across the jet width at all observed frequencies, indicative of a gradient in the emission-weighted intrinsic polarization angle across the jet and possibly a systematic gradient in the rotation measure; moreover, the percentage of polarization increases near the curvature at C, likely consistent with the presence of a local jet-ISM interaction and/or changing magnetic-field directions. The hot spot B shows a higher rotation measure, and has no detected proper motion. These facts provide some evidence for a stationary shock in the vicinity of B. Comparison of the present VLBI observations with those made 8.43 years ago suggests a significant northward motion for A2 with an apparent transverse velocity $\\beta_{app}=3.7\\pm0.4\\, c$. The apparent superluminal motion suggests that the relativistic jet plasma moves at a velocity of $\\gtrsim0.96\\, c$ if the jet is viewed at an inclination angle less than $20\\degr$. A simple precessing jet model and a hydrodynamical isothermal jet model with helical-mode Kelvin-Helmholtz instabilities are used to fit the oscillatory jet trajectory of 3C~48 defined by the bright knots. ",
        "introduction": "\\label{section:intro}  Compact Steep Spectrum (CSS) sources are a population of powerful radio sources with projected linear size less than 20 kpc and steep high radio frequency spectrum $\\alpha<-0.5$ \\footnote{In the present paper, the spectral index is defined as $S_\\nu\\propto\\nu^\\alpha$.} (Peacock \\& Wall 1982, Fanti et al. 1990, and review by O'Dea 1998 {\\bf and Fanti 2009}). Kinematical studies of the hot spots and analysis of the high-frequency turnover in the radio spectrum due to radiative cooling imply ages for CSS sources in the range 10$^2$--10$^5$ yr (e.g., Owsianik, Conway \\& Polatidis 1998; Murgia et al. 1999). The sub-galactic size of CSS sources has been used to argue that CSS sources are probably young radio sources, (the `youth' model: Fanti et al. 1995; Readhead et al. 1996). However, another interpretation attributes the apparent compactness of the CSS sources to being strongly confined by the dense ISM in the host galaxy (the `frustration' model: van Breugel, Miley \\& Heckman 1984). Spectroscopic observations of CSS sources provide evidence for abundant gas reservoirs in the host galaxies and strong interaction between the radio sources and the emission-line clouds \\cite{ODe02}. Some CSS sources have been observed to have high-velocity clouds (as high as $\\sim500$ km\\,s$^{-1}$) in the Narrow-Line Region (NLR), presumably driven by radio jets or outflows; an example is 3C~48 \\cite{Cha99,Sto07}. In addition, many CSS sources show distorted radio structures, suggestive of violent interaction between the jet and the ambient interstellar medium \\cite{Wil84,Fan85,Nan91,Spe91,Nan92,Aku91}. The ample supply of cold gas in their host galaxies and their strong radio activity, which results in a detection rate as high as $\\sim30$ per cent in flux-density limited radio source surveys \\cite{Pea82,Fan90}, make CSS sources good laboratories for the study of AGN triggering and feedback.  3C~48 ($z=0.367$) is associated with the first quasar to be discovered \\cite{Mat61,Gre63} in the optical band. Its host galaxy is brighter than that of most other low redshift quasars. The radio source 3C~48 is classified as a CSS source due to its small size and steep radio spectrum \\cite{Pea82}. Optical and NIR spectroscopic observations suggest that the active nucleus is located in a gas-rich environment and that the line-emitting gas clouds are interacting with the jet material \\cite{Can90,Sto91,Cha99,Zut04,Kri05,Sto07}. VLBI images \\cite{Wil90,Wil91,Nan91,Wor04} have revealed a disrupted jet in 3C~48, indicative of strong interactions between the jet flow and the dense clouds in the host galaxy. Although some authors \\cite{Wil91,Gup05,Sto07} have suggested that the vigorous radio jet is powerful enough to drive massive clouds in the NLR at speeds up to 1000 km~s$^{-1}$, the dynamics of the 3C~48 radio jet have yet to be well constrained. Due to the complex structure of the source, kinematical analysis of 3C~48 through tracing proper motions of compact jet components can only be done with VLBI observations at 4.8 GHz and higher frequencies, but until now the required multi-epoch high-frequency VLBI observations had not been carried out.  In order to study the kinematics of the radio jet for comparison with the physical properties of the host galaxy, we observed 3C~48 in full polarization mode with the VLBA at 1.5, 4.8 and 8.3 GHz in 2004, and with the EVN and MERLIN at 1.65 GHz in 2005. Combined with earlier VLBA and EVN observations, these data allow us to constrain the dynamics of the jet on various scales. Our new observations and our interpretation of the data are presented in this paper. The remainder of the paper is laid out as follows. Section 2 describes the observations and data reduction; Section 3 presents the total intensity images of 3C~48; and Section 4 discusses the spectral properties and the linear polarization of the components of the radio jet. In Section 5, we discuss the implications of our observations for the kinematics and dynamics of the radio jet. Section 6 summarizes our results. Throughout this paper we adopt a cosmological model with Hubble constant $H_0$=70 km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_m=0.3$, and $\\Omega_{\\Lambda}=0.7$. Under this cosmological model, a 1-arcsec angular separation corresponds to a projected linear size of 5.1 kpc in the source frame at the distance of 3C 48 ($z=0.367$). ",
        "conclusions": "We have observed 3C~48 at multiple frequencies with the VLBA, EVN and MERLIN with spatial resolutions between tens and hundreds of parsec. Our principal results may be summarized as follows:  (1) The total-intensity MERLIN image of 3C~48 is characterized by two components with comparable integrated flux density. A compact component aligns with the VLBI jet, while an extended envelope surrounds it. The extended emission structure becomes diffuse and extends toward the northeast at $\\sim$0.25 arcsec from the nucleus. The extended component shows a steeper spectrum than the compact jet.  (2) In the VLBA and EVN images, the compact jet seen in the MERLIN image is resolved into a series of bright knots. Knot A is further resolved into two smaller features A1 and A2 in 4.8- and 8.3-GHz VLBA images. A1 shows a flat spectrum with spectral index $\\alpha^{4.8}_{8.3}=-0.34\\pm0.04$. A2 shows a steep spectrum with $\\alpha^{4.8}_{8.3}=-1.29\\pm0.16$, and may be identified with the inner jet. The brightness temperature of A1 is $>10^9$~K and much higher than the $T_b$ of A2. The flux densities of A1 and A2 in epoch 2004 show a 100 and 60 per cent decrease compared with those in 1996. The high brightness temperature, flat spectrum and variability imply that A1 is the synchrotron self-absorbed core found close to the active nucleus.  (3) Comparison of the present VLBA data with those of 1996 January 20 strongly suggests that A2 is moving, with an apparent velocity $3.7c\\pm0.4c$ to the North. Combining the apparent superluminal motion and the jet-to-counterjet intensity ratio yields a constraint on the jet kinematics and geometry: the jet is relativistic ($>0.85c$) and closely aligned to the line of sight ($<35\\degr$).  (4) We present for the first time VLBI polarization images of 3C~48, which reveal polarized structures with multiple sub-components in component C. The fractional polarization peaks at the interface between the compact jet and the surrounding medium, perhaps consistent with a local jet-induced shock. The systematic gradient of the EVPAs across the jet width at C can be attributed to the combination of a gradient in the emission-weighted intrinsic polarization angle across the jet and possibly a systematic gradient in the RM. Changing magnetic field directions are a possible interpretation of the RM gradient, but other alternatives can not be ruled out. The fractional polarization of the hot spot B increases towards higher frequencies, from $\\sim1$ per cent (1.6 GHz), $\\sim2.0$ per cent (4.8 GHz) to $12$ per cent (8.3 GHz). The relatively low degree of polarization at lower frequencies probably results from a unresolved Faraday screen associated with the NLR clouds and/or the internal depolarization in the jet itself. Hot spot B has a higher RM than C, which can perhaps be attributed to a stationary shock in the vicinity of B. The core A at all frequencies is unpolarized, which may be the result of a tangled magnetic field in the inner part of the jet.  (5) The combined EVN+MERLIN 1.65-GHz image and 1.5-GHz VLBA images show that the bright knots trace out a wave-like shape within the jet. We fitted the jet structure with a simple precession model and a K-H instability model. Both models in general reproduce the observed oscillatory jet trajectory, but neither of them is able to explain all the observations. More observations are required to investigate the physical origin of the helical pattern. Further monitoring of the proper motion of the inner jet A2 should be able to constrain the ballistic motion in the framework of the precessing jet. High-resolution VLBI images of the inner jet region will be required to check whether or not the jet flow is oscillating on scales of tens of mas, which might give a morphological means of discriminating between the two models. Sophisticated simulations of the jet would be needed to take into account the deceleration of the jet flow due to kinetic energy loss via jet-cloud interaction and radiation loss, but these are beyond the scope of the present paper."
    },
    "0910/0910.4394_arXiv.txt": {
        "abstract": "In order to explain the inflated radii of some transiting extrasolar giant planets, we investigate a tidal heating scenario for the inflated planets WASP-4b, WASP-6b, WASP-12b, WASP-15b, and TrES-4.  To do so, we assume that they retain a nonzero eccentricity, possibly by dint of continuing interaction with a third body. We calculate the amount of extra heating in the envelope that is then required to fit the radius of each planet, and we explore how this additional power depends on the planetary atmospheric opacity and on the mass of a heavy-element central core. There is a degeneracy between the core mass $M_{\\rm core}$ and the heating $\\dot{E}_{\\rm heating}$. Therefore, in the case of tidal heating, there is for each planet a range of the couple $\\{M_{\\rm core},e^2/Q'_p\\}$ that can lead to the same radius, where $Q'_{p}$ is the tidal dissipation factor and $e$ is the eccentricity. With this in mind, we also investigate the case of the non-inflated planet HAT-P-12b, which can admit solutions combining a heavy-element core and tidal heating.  A substantial improvement of the measured eccentricities of such planetary systems could simplify this degeneracy by linking the two unknown parameters $\\{M_{\\rm core},Q'_{p}\\}$. Further independent constraints on either of these parameters would, through our calculations, constrain the other. ",
        "introduction": "\\label{sec:intro} The transiting extrasolar giant planets (EGPs), around 60 discovered to date\\footnote{See J. Schneider's Extrasolar Planet Encyclopaedia at http://exoplanet.eu, the Geneva Search Programme at http://exoplanets.eu, and the Carnegie/California compilation at http://exoplanets.org.}, offer the best testbed for theoretical models of the evolution of planetary radii.  In the last 14 years, many such theoretical models and tests have been proposed and investigated \\citep{guillot_et_al1996, burrows_et_al2000, Bodenheimer_et_al_2001, burrows_et_al2003, Baraffe_et_al_2003, Bodenheimer_et_al_2003, Gu_et_al_2003,burrows_et_al2004, Fortney_and_Hubbard_2004, Baraffe_et_al_2004, Chabrier_et_al_2004, laughlin_et_al_2005_1, Baraffe_et_al_2005, Baraffe_et_al_2006, burrows_et_al2007, Fortney_et_al_2007, Marley_et_al_2007, chabrier+baraffe2007, Liu_et_al_2008, Baraffe_et_al_2008, Ibgui_and_Burrows_2009, Miller_et_al_2009, Leconte_et_al_2009}.  A gas giant planet's radius is a function of many variables, including the planet's mass and age; the stellar irradiation flux and spectrum; the composition -- in particular, the heavy-element content -- of the atmosphere, the envelope, and the core; the atmospheric circulation that couples the day and the night sides; and any processes that could generate an extra power source in the interior of the planet, such as tidal heating.  Furthermore, the connection between a planet's physical radius and its transit radius is complicated by the transit radius effect \\citep{burrows_et_al2003, Baraffe_et_al_2003}.  Since each of these variables can change from one planet-star system to another, only a custom calculation can determine if the measured transit radius of a giant planet matches theoretical predictions.  The first observations of a transiting planet (HD209458b), which found that its diameter is more than 30\\% greater than Jupiter's \\citep{henry_et_al2000, charbonneau_et_al2000}, revealed a gap in our understanding of radius evolution. HD209458b's radius is significantly larger than the standard evolutionary theory would predict for an object of its age \\citep{knutson_et_al2007a, burrows_et_al2007}.  For this reason it is said to be inflated, or bloated.  In the last decade, a number of other inflated planets have been discovered, whose radii are similarly difficult to explain in terms of simple evolutionary theory (CoRoT-1b: \\citealt{Barge_et_al_2008, Gillon_et_al_2009_3}, CoRoT-2b: \\citealt{Alonso_et_al_2008}, HAT-P-1b: \\citealt{bakos_et_al2007a, Winn_et_al_2007, Johnson_et_al_2008}, TrES-2: \\citealt{odonovan_et_al2006b}, TrES-4: \\citealt{Mandushev_et_al_2007, Sozzetti_et_al_2008}, WASP-4b: \\citealt{Wilson_et_al_2008, Southworth_et_al_2009_2, Gillon_et_al_2009_1, Winn_et_al_2009_1}, WASP-6b: \\citealt{Gillon_et_al_2009_2}, WASP-12b: \\citealt{Hebb_et_al_2009}, WASP-15b: \\citealt{West_et_al_2009}, XO-3b: \\citealt{Johns-Krull_et_al_2008, Winn_et_al_2008}).  The present paper investigates a stationary heating scenario as a possible explanation of the inflated transiting giant planets recently discovered: WASP-4b, WASP-6b, WASP-12b, and WASP-15b.  In addition, we revise the case of TrES-4, already discussed by \\citet{Liu_et_al_2008}, and consider HAT-P-12b \\citep{Hartman_et_al_2009}, a planet that is not inflated, but that appears to require a massive, dense core to explain its small transit radius. At the time of the writing of this paper, three other inflated transiting EGPs have been discovered: CoRoT-5b \\citep{Rauer_et_al_2009}, HAT-P-13b \\citep{Bakos_et_al_2009_2}, and WASP-17b \\citep{Anderson_et_al_2009}.  Although we might consider them in detail in a future study, we find it reassuring that preliminary investigations of these objects indicate that the general philosophy of the present paper can be applied to them as well.  Tidal heating as an explanation for inflated close-in EGPs was proposed by \\citet{Bodenheimer_et_al_2001, Bodenheimer_et_al_2003}. This heating being due to a nonzero eccentricity, they suggested the necessity of an excitation mechanism, as for example gravitational interaction with another planetary companion.  This possibility was examined by \\citet{Mardling_2007}.  \\citet{Liu_et_al_2008} investigated the cases of the inflated planets TrES-4, XO-3b, and HAT-P-1b, while assuming a constant rate of interior heating. \\citet{Ibgui_and_Burrows_2009}, \\citet{Ibgui_et_al_2009_2}, and \\citet{Miller_et_al_2009} coupled the evolution of the planetary radius with that of the orbit, under the influence of tidal interactions.  The interesting feature revealed by this scenario is a possible transient inflation of the planetary radius that temporarily interrupts its monotonic standard shrinking.  Here, we assume a constant heating rate. Such an assumption would be valid, in the context of tidal heating, if there is a stationary scenario such that the orbital eccentricity and semimajor axis of the transiting EGP are quasi-constant and the radius has reached a ``quasi-equilibrium'' value.  This paper is organized as follows. In Section~\\ref{sec:properties}, we present the relevant properties of the planet-star systems in our study.  In Section~\\ref{sec:model}, we briefly summarize the model assumptions adopted for our planetary radius evolutionary calculations.  In Section~\\ref{sec:theo_Req}, we present both the details of our extra-heating model and its results. We derive the necessary heating rates $\\dot{E}_{\\rm heating}$ for the equilibrium radii $R_{\\rm eq}$ to match the measured values. In Section \\ref{subsec:no_central_core}, we assume that the planets have no heavy-element central core and show that an enhanced planet atmospheric opacity, since this results in a larger radius, requires less internal heating.  In Section \\ref{subsec:central_core_solar}, we freeze the planetary atmospheric opacity, and show that the presence of a heavy-element central core, whose effect is to shrink the planetary radius, requires more internal heating.  In Sections \\ref{subsec:no_central_core} and \\ref{subsec:central_core_solar}, we relate $\\dot{E}_{\\rm heating}$ to the ratio $e^2/Q'_p$, where $Q'_p$ is the tidal dissipation factor \\citep{Goldreich_1963} and $e$ the orbital eccentricity.  In Section \\ref{subsec:compatible_models} we further show the following degeneracy: a given radius can be matched by a range of values of the couples $\\{M_{\\rm core},\\dot{E}_{\\rm heating}\\}$ or $\\{M_{\\rm core},e^2/Q'_p\\}$. When the measured eccentricities become sufficiently reliable, we could directly relate $M_{\\rm core}$ and $Q'_p$. We make a first attempt at such a relation in this paper, while we remain fully aware of the poor constraints on $e$.  We conclude and provide a general discussion in Section \\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion} We have proposed in this paper a stationary tidal heating scenario to explain the radii of the inflated transiting EGPs WASP-4b, WASP-6b, WASP-12b, WASP-15b, and TrES-4. A constant heating rate may be achieved with quasi-constant eccentricity and semimajor axis, and after the planetary radius has reached its equilibrium value. Such a scenario might explain the inflated radii of some transiting EGPs. We have calculated the amount of additional interior heating that is required to fit the radius of each planet, and explored how it depends on the planet's atmospheric opacity and on a heavy-element central core. There is a degeneracy between the required heating rate to fit the radius, the atmospheric opacity, and the heavy-element core mass inside the planet. In terms of tidal heating, there is, for a given opacity, a locus of points $\\{M_{\\rm core},e^2/Q'_p\\}$ that can lead to the same radius. For this reason, a substantial improvement in the precision of measured orbital eccentricities would transform this degeneracy into a reliable $\\{M_{\\rm core},Q'_p\\}$ degeneracy. Combined with theories of tidal dissipation inside the planets that constrain the values of $Q'_p$, and with planet formation models that may also constrain the values of $M_{\\rm core}$, this scenario could represent an explanation of these inflated radii based on tidal heating.  Our general conclusions are the following: \\begin{itemize}\\itemsep0cm \\item The heating rate, $\\dot{E}_{\\rm heating}$, necessary to fit an observed inflated radius is generally small in comparison to the irradiation rate, $\\dot{E}_{\\rm irradiation}$, with a ratio $\\dot{E}_{\\rm heating}/\\dot{E}_{\\rm irradiation}$ ranging from $10^{-6}$ to $10^{-1}$. \\item The higher the heating rate, the larger the equilibrium radius. On the other hand, the larger the core mass, the smaller the equilibrium radius, and the higher the required heating rate in order to fit the measured radius. In other words, $\\dot{E}_{\\rm heating}$ increases with $M_{\\rm core}$. \\item Tidal heating might provide the extra interior power that inflates some planets.  If tides are the extra power source, then the tidal dissipation factor $Q'_p$ is a decreasing function of core mass $M_{\\rm core}$, but an increasing function of the atmospheric opacity. It may also be different for different planets. \\item The more massive the planet, the less its radius is sensitive to extra heating. \\item If the steady-state scenario presented in this paper obtains in a planetary system, the presence and the characterization of a central heavy-element core in that planet can be constrained by a better knowledge of the eccentricity, the planet's atmospheric opacity, and the tidal dissipation in its interior governed by $Q'_p$. \\item It would be interesting to revisit the work of \\citet{Mardling_2007} in the specific instances of the planets under consideration in this paper, so as to characterize what possible companions might cause quasi-constant eccentricity and semimajor axis. Better constraints from radial velocity data will reveal which such companions, if any, might be present. \\end{itemize}  The applicability of the steady-state scenario should be examined for each planet individually. Assuming the current best measured estimates of the eccentricities, and being fully aware of the poor knowledge of this parameter and, therefore, that the following comments will need to be revised when more accurate measurements are available, we can state more specifically, planet by planet, that: \\begin{itemize}\\itemsep0cm \\item TrES-4, for which we assume $e$=0.01, is the most difficult to reconcile with the steady-state scenario. The tidal dissipation parameter $Q'_p$ might be $10^{3.5}$ for a solar atmospheric opacity, lower than theoretical models or empirical determinations would generally suggest.  If, as is likely, the actual eccentricity is lower than 0.01, the required value of $Q'_p$, in the context of our model, would be even lower.  On the other hand, if the atmospheric opacity is greater than solar, the planet would cool more slowly and the required value of $Q'_p$, in our model, would be greater (for instance, $10^{4.7}$ for a $10\\times$solar opacity). For the purpose of constraining evolutionary models of this planet, it would be useful to study in greater depth both its eccentricity, which is not well constrained, and the composition and opacity of its atmosphere. \\item WASP-4b can be fit, but with values of $Q'_p$ or $M_{\\rm core}$ that are higher than we would expect. Its actual eccentricity might be noticeably lower than the upper limit currently available. \\item WASP-6b can be fit within $1\\sigma$ with a $3\\times$solar opacity, without any extra heating. Consequently, a lower opacity involves an extra heating. \\item WASP-12b requires either high values of $Q'_p$ or extremely high core masses, such as 300~$M_{\\earth}$. However, such a high $M_{\\rm core}$ might be ruled out by planetary formation models. \\item WASP-15b can be fit with $Q'_p$ ranging in the expected interval $10^{5}-10^{6}$ and with core masses below $\\sim$$100 M_{\\earth}$. \\item HAT-P-12b, the only non-inflated planet, can be fit without invoking tidal heating, and with a core mass smaller than $\\sim$$50 M_{\\earth}$. However, models involving tidal heating and a central core are also compatible with its observed properties, within the steady-state assumption of this paper. \\end{itemize}  There are several caveats that the reader should bear in mind. Details of tidal dissipation theory remain to be determined, specifically the mechanism and location of heat deposition (in the radiative mantle or in the convective envelope).  Also, we need to better understand the difference between day and night cooling.  This may require full three dimensional general circulation models with radiative transfer.  Finally, we have assumed that the core mass and the atmospheric opacity are independent.  Broadly speaking, it seems likely that these properties are somewhat connected, due to a probable correlation of the compositions of the central core, the convective envelope, and the atmosphere. However, the lack of a unique mapping between core mass and atmospheric opacity might limit the extent to which future investigations might be able to treat these variables as coupled.  The fact that some transiting extrasolar giant planets have inflated radii continues to pose a theoretical challenge.  The stationary tidal heating scenario described in this paper might explain the puzzle in some cases.  Further observational constraints will clarify the viability of this scenario.  To conclude, we mention the possibility of an alternative model, that might also be able to explain the inflated radii of some of the planets under consideration in this paper -- namely, the coupled evolution of the planetary radius and its orbit, in the case of a two-body gravitational and tidal interaction \\citep{Ibgui_and_Burrows_2009, Ibgui_et_al_2009_2, Miller_et_al_2009}. At this stage, both scenarios appear to be plausible."
    },
    "0910/0910.4488_arXiv.txt": {
        "abstract": "We present spatially-resolved dynamics for six strongly lensed star-forming galaxies at $z = 1.7$--3.1, each enlarged by a linear magnification factor $\\sim\\times$8.  Using the Keck laser guide star AO system and the OSIRIS integral field unit spectrograph we resolve kinematic and morphological detail in our sample with an unprecedented fidelity, in some cases achieving spatial resolutions of $\\simeq$100\\,pc.  With one exception our sources have diameters ranging from 1--7 kpc, integrated star formation rates of $2-40\\,\\Msolyr$ (uncorrected for extinction) and dynamical masses of $10^{9.7-10.3} \\Msol$. With this exquisite resolution we find that four of the six galaxies display coherent velocity fields consistent with a simple rotating disk model. Our model fits imply ratios for the systemic to random motion, $V_c \\sin{i} / \\sigma$, ranging from 0.5--1.3 and Toomre disk parameters $Q < 1$.  The large fraction of well-ordered velocity fields in our sample is consistent with data analyzed for larger, more luminous sources at this redshift. We demonstrate that the apparent contradiction with earlier dynamical results published for unlensed compact sources arises from the considerably improved spatial resolution and sampling uniquely provided by the combination of adaptive optics and strong gravitational lensing.  Our high resolution data further reveal that all six galaxies contain multiple giant star-forming H{\\sc ii} regions whose resolved diameters are in the range 300 pc -- 1.0 kpc, consistent with the Jeans length expected in the case of dispersion support. From the kinematic data we calculate that these regions have dynamical masses of $10^{8.8-9.5} \\Msol$, also in agreement with local data. However, the density of star formation in these regions is $\\sim$100$\\times$ higher than observed in local spirals; such high values are only seen in the most luminous local starbursts. The global dynamics and demographics of star formation in these H{\\sc ii} regions suggest that vigorous star formation is primarily governed by gravitational instability in primitive rotating disks. The physical insight provided by the combination of adaptive optics and gravitational lensing suggests it will be highly valuable to locate many more strongly-lensed distant galaxies with high star formation rates before the era of the next-generation ground-based telescopes when such observations will become routine. ",
        "introduction": "\\label{secintro} Studies of star-forming galaxies at high redshift ($z > 1.5$) over the past decade have mapped the demographics of the populations as a whole yielding valuable data on their integrated star formation rates, stellar and dynamical masses, metallicities and morphologies \\citep{Shapley01,Shapley03,Reddy04,Erb06, Law07a}. From these data, a consistent picture of the development of the comoving density of star formation and its contribution to the present day stellar density has emerged \\citep{Dickinson03, Hopkins06}. Following a rapid rise in activity at early times, corresponding to the redshift range $2<z<6$, the star formation rate has decline markedly over the past 8 Gyr corresponding to the interval $0<z<1$. Of particular interest is the relationship between sources observed at the time close to the peak epoch of star formation, $z = 2-3$ and the population of massive galaxies observed today.  Recent theoretical work has focused on the mechanisms via which early galaxies assemble their stars and evolve morphologically to the present day Hubble sequence. Minor mergers and cold stream gas accretion have emerged as a possible means of building disks, central bulges and elliptical galaxies from early star forming Lyman break galaxies \\citep{Dekel09, Brooks09}. These models of galaxy evolution rely on simple descriptions of complex physical processes such as gas cooling, star formation and feedback mechanisms.  It is therefore crucial to undertake relevant observations to constrain these processes.  The most direct observational route to making progress is high quality resolved spectro-imaging of early galaxies which can be used to determine both their dynamical state and the distribution of star-formation.  Integral field unit (IFU) spectroscopy of nebular emission lines is now yielding valuable kinematic data for a growing sample of luminous high-redshift star forming galaxies \\citep{Genzel06, Genzel08, Forster06, Forster09, Law07, Law09}. Although the galaxies observed so far do not yet in any way constitute a complete well-defined sample, some important results have emerged. The data reveal a mix of dispersion-dominated systems, rotating systems, and major mergers with the common observation that all galaxies studied so far show relatively high velocity dispersions of $\\sim 50-100$ km\\,s$^{-1}$. A larger fraction of the more massive ($M_{dyn} \\gtrsim 10^{11} M_{\\odot}$) galaxies appear to be rotationally supported, with lower mass galaxies ($M_{dyn} \\lsim 10^{10}M_{\\odot}$) tending to have dispersion-dominated kinematics \\citep{Law09}.  However, these studies are hampered by the poor spatial resolution inherent in studies of distant star-forming galaxies whose typical sizes are $\\simeq1$--5\\,kpc. Even with adaptive optics on a 8--10 meter aperture, the physical resolution is limited to FWHM $\\gtrsim 1$ kpc which means only a few independent resolution elements. Consequently it is unclear whether the velocity shear observed in the smaller 2--5\\,kpc sources arises from rotation or merging. Moreover, although there are suggestions that distant star-forming regions have sizes much larger than those of local H{\\sc ii} regions \\citep{Elmegreen05}, the claim remains uncertain since the distant regions are not properly resolved.  Thus, the diffraction limit of the current generation of optical/infrared telescopes limits progress, even with all the adaptive optics tools of the trade.  Fortunately, sources which undergo strong gravitational lensing are enlarged as seen by the observer and can thus be studied at much higher physical resolution in their source plane \\citep{Nesvadba06, Swinbank07}. This improvement in physical resolution enables us to better distinguish various forms of velocity field and to examine the properties of star-forming giant H{\\sc ii} regions.  A convincing demonstration of the benefits of lensing as applied to this topic was the work of \\cite{Stark08} which analyzed IFU data on a $z=3.07$ galaxy magnified in area by a factor of 28$\\times$ demonstrating resolved dynamics on $\\simeq$100 pc scales. This represents an order of magnitude improved sampling compared to the earlier studies cited above.  Here we present IFU observations of a further five strongly lensed galaxies which we add to the source studied by \\cite{Stark08}. In addition to studying the well-sampled kinematic structure of a population of sub-L$^{\\ast}$ star-forming galaxies at $z\\simeq2$, our fine resolution allows us to examine the size, dynamical mass, and luminosity of typical star-forming regions, thereby probing the conditions in which most of the star formation takes place. We find that all of our galaxies contain multiple, distinct star forming regions.  These results offer a preview of the science which will be possible with 30 meter class optical/near-IR telescopes with a diffraction limit comparable to the resolution of our data; they provide the most detailed view to date of typical high-redshift star forming galaxies.  A plan of the paper follows. In $\\S$2 we discuss our selection of targets and Keck observations, including the production and use of gravitational lens modeling essential for achieving a high resolution in the source plane. In $\\S$3 we analyze our spectroscopic results, discussing each object in turn and summarizing the kinematic data. In $\\S$4 we examine the physical scale of our star forming regions in comparison to those seen locally. We summarize our conclusions in $\\S$5.  Throughout we use a {\\it WMAP} cosmology \\citep{Spergel03} with $\\Omega_{\\Lambda}$=0.73, $\\Omega_{m}$=0.27, and H$_{0}$=72\\,km\\,s$^{-1}$\\,Mpc$^{-1}$.  All quoted magnitudes are in the AB system unless otherwise noted. ",
        "conclusions": "We have utilized strong gravitational lensing together with laser-assisted guide star adaptive optics at the Keck observatory to study the internal characteristics of six $z = 1.7-3$ star-forming galaxies with a spatial resolution down to $\\sim 100-200$ parsecs. Extending beyond the diffraction limited capability of current 8--10 meter class telescopes, our study provides an interesting preview of the science that will routinely be possible with future 30 meter class optical/infrared telescopes with adaptive optics systems (Ellis 2009). We demonstrate that enhanced resolution allows us to resolve morphological and kinematic structure which has not been (and cannot be) discerned in AO-based surveys of similar non-lensed sources. In particular we resolve multiple giant star forming regions of size consistent with that expected from Toomre instability in {\\sl all} of our targets (including separate merging components). Furthermore we determine the dynamical state of galaxies which would be too poorly sampled to distinguish between rotation, merging, or dispersion-dominated kinematics without gravitational lensing.  The lensed sample allows a unique study of relatively faint and small galaxies at high resolution. The median luminosity of the sample is half of the characteristic $L^*$ for $z=3$ LBGs and well below that of other high-redshift IFU studies. We find that the kinematics of the sub-$L^*$ population are in general agreement with larger and more luminous galaxies at similar redshifts: four of the six systems are clearly rotating with $V\\sin{i}/\\sigma = 0.5$--1.3, one has a velocity gradient consistent with rotation within the small H$\\alpha$-bright region of the arc detected with OSIRIS, and one is a major merger. All have high velocity dispersion $\\sigma = 50$--$100\\kms$, consistent with all other resolved observations of $z>2$ star-forming galaxies and comparable to the value $\\sigma \\simeq 50 \\kms$ for stars which formed in the Galactic disk at $z\\simeq2$ \\citep{Sparke00}. Galaxies with a larger radius of detected nebular emission tend to have higher velocity shear and $V/\\sigma$ and lower $\\Sigma_{SFR}$; these trends are consistent with other observations. The case of MACS\\,J0744+3927 with inclination-corrected $V=180^{+100}_{-40} \\kms$ clearly demonstrates that some small ($R=1.0\\pm0.2$\\,kpc), fast-rotating field disk galaxies are already in place by $z=2$. Well-sampled velocity fields demonstrate that the ubiquitous giant clumps are often embedded in common rotating systems and are not independent merging components. In contrast, non-lensed clumpy systems with velocity shear have been interpreted as mergers or remained ambiguous. These results demonstrate the need of high-resolution observations to distinguish rotating systems and mergers, as well as to probe the $< 1$\\,kpc scale of star formation at high redshift.  The observations presented here represent the best current probe of the scale of star formation at $z \\gtrsim 2$ and, as such, can be compared to theoretical predictions. Most of the star formation in our sample occurs within giant H{\\sc ii} regions of diameter 0.3--1 kpc comparable to the largest local H{\\sc ii} regions, with star formation rate density comparable to the most vigorous local starbursts and $\\sim 100 \\times$ higher than in typical spiral galaxies. This offset in $\\Sigma_{SFR}$ cannot be explained by the different resolution or sensitivity of low- and high-redshift observations. Only one of our six objects is a clear major merger with both merging components broken into high-SFR clumps, suggesting that star formation episodes in sub-L$^*$ high redshift galaxies are triggered by Toomre instability independent of major merger events. We note that the merging system demonstrates no enhancement in star formation (as traced by nebular emission) compared to the non-merging galaxies.  The increased high-mass star formation relative to local spirals likely results from some combination of higher gas density, increased star formation efficiency, shorter star formation timescales, and possibly other effects. The clumpy star formation, high $\\Sigma_{SFR}$, kinematics consistent with rotation, and small $Q<1$ are in qualitative agreement with the galaxy evolution model of \\cite{Dekel09} in which galaxies accrete their baryonic mass from cold streams. In this model the clumps form due to Toomre instability in turbulent disks and migrate into the galaxy center on timescales of a few hundred Myr, forming a bulge which stabilizes the system and increases the star formation timescale after $\\sim 2$ Gyr. Numerical simulations suggest that cold-stream accretion is a dominant mechanism in galaxy assembly, especially at early times and in galaxies with halo masses less than a few $10^{11} \\Msol$ \\citep{Keres09,Ocvirk08,Brooks09}. In particular the high-resolution simulations of \\cite{Brooks09} show that accretion of cold gas dominates the buildup of stellar disks and bulges in sub-$L^*$ galaxies and allows disk formation at earlier times than if the accreted gas is shock-heated. The rotating kinematics and distribution of star formation in the lensed sample therefore supports, at least qualitatively, a cold-stream accretion scenario for galaxy formation.  Analysis of our sample has shown the benefit of using lensed sources to probe the inner structure of early star-forming galaxies. The advent of many panoramic multi-color imaging surveys will hopefully reveal many further examples of such systems, ensuring further progress with laser-assisted guide star adaptive optics. Ultimately, the next generation of 30-meter optical/near-infrared telescopes will extend the high-resolution studies presented herein to the entire population of high-redshift star forming galaxies."
    },
    "0910/0910.1925_arXiv.txt": {
        "abstract": "{We present a study of the mixing properties of the simulated intra cluster medium, using tracers particles that are advected by the gas flow during the evolution of cosmic structures. Using a sample of seven galaxy clusters (with masses in the range of $M \\sim 2 - 3 \\cdot 10^{14}M_{\\odot}/h$) simulated with a peak resolution of $25kpc/h$ up to the distance of two virial radii from their centers, we investigate the application of tracers to some important problems concerning the mixing of the ICM. The transport properties of the evolving ICM are studied through the analysis of pair dispersion statistics and mixing distributions. As an application, we focus on the transport of metals in the ICM. We adopt simple scenarios for the injection of metal tracers in the ICM, and find remarkable differences of metallicity profiles in relaxed and merger systems also through the analysis of simulated emission from Doppler-shifted Fe XXIII lines.} ",
        "introduction": "\\label{sec:intro}  The mixing properties of the intra cluster medium (ICM) are still poorly understood and a number of open fields of research are tightly connected to this important topic: the 'heating cooling flows' problem (e.g. Bruggen \\& Kaiser 2002; Omma et al. 2004), the spreading of metals in the ICM (e.g. Schindler \\& Diaferio 2008; Borgani et al.2008), the stability of thermal profiles in the innermost region of galaxy clusters (e.g. Sharma et al.2009), the efficiency in the ram pressure stripping of accreting satellites and the emergence of cold fronts in galaxy clusters (e.g. Ascasibar \\& Markevitch 2006; Markevitch \\& Vikhlinin 2007) and many others.  From the theoretical point of view, in the last few years meaningful progress has been made in understanding the convective properties of the ICM. Focusing on the role played by instabilities in a magnetized, low-collisional plasma (e.g. Parrish \\& Stone 2005; Quataert 2008; Ruszkowski \\& Oh 2009), a new picture of the ICM stability was presented, which drastically alters the ``classic'' picture provided by the Schwarzschild criterion, in which stability is ensured by $dlog(S)/dr>0$ (Schwarzschild 1959), where $S$ is the specific entropy.  These approaches cannot take into account all sources of mixing motions in galaxy clusters though, which are due to large scale accretions of matter and turbulent motions. Indeed there is clear evidence nowadays that a sizable amount of turbulent motions may be present in the ICM. Observations suggest large scale subsonic turbulent motions in the range of $\\sim 100\\rm kpc-1 \\rm Mpc$ (e.g. Schuecker et al.2004; Henry et al.2004; Churazov et al.2004). In addition, also studies of Faraday Rotation allow a complementary approach and suggest that the ICM magnetic field is  tangled over a broad range of scales (e.g. Murgia et al.2004; Vogt \\& Ensslin 2005; Kuchar \\& Ensslin 2009; Bonafede et al.2010); also, the diffuse radio emission detected in a fraction of merging clusters may provide indirect evidence for turbulence in the ICM (e.g. Brunetti et al. 2008).  Obtaining a self-consistent picture of the evolving ICM, where large scale and small scale mixing properties of galaxy clusters are followed during the whole cluster evolution is challenging, and in this respect cosmological high resolution numerical simulations may provide additional and valuable information.  At present, two main numerical methods are massively applied to cosmological numerical simulations: Lagrangian methods, which sample both the Dark Matter and the gas properties using unstructured point-like fluid elements, usually regarded as particles (e.g. smoothed particles hydrodynamic codes such as GADGET, Springel 2005, or HYDRA, Couchman et al.1995) and Eulerian methods, which reconstruct the gas properties with a discrete sampling of the space using meshes of various resolution (e.g. cosmological TVD code, Ryu et al.1993; ENZO, Norman et al.2007; FLASH, Fryxell et al.2000; RAMSES, Teyssier 2002 etc) and model the Dark Matter properties with a particle mesh approach.   The typical advantages and drawbacks of the two approaches have been extensively investigated. Lagrangian methods allow one to obtain a high spatial resolution where matter concentrates, but in low density regions the sampling is rather poor. By construction, this class of methods allows one to follow in a natural way the advection of a single parcel of (gas or DM) matter during the whole simulation and its dynamical evolution.  Eulerian methods have a spatial resolution fixed to the size of the cell in the computational mesh, and this often is inadequate to follow the details of cosmological and physical processes of interest in the innermost region of collapsed objects. The adaptive mesh refinement technique (AMR) overcomes this problem by increasing the spatial resolution of the Eulerian mesh in regions of particular interest (e.g. Berger \\& Oliger 1984; Berger \\& Colella 1989; Kravtsov et al.1997). A further important feature of Eulerian methods is that they are strictly conservative for the total fluid energy, and therefore they accurately follow the Rankine-Hugoniot relations for shocks (a property shared by shock capturing methods, also known as Godunov schemes, Godunov 1959). This is particular important in cosmological simulations, where cosmic shocks play an important role in setting both thermal and non thermal properties of the ICM (e.g. Ryu et al.2003; Pfrommer et al.2006; Vazza et al.2009).  In the last few years a number of works highlighted and investigated the causes which lead to the main differences between these numerical approaches (Agertz et al.2006; Tasker et al.2008; Wadsley et al. 2008; Mitchell et al.2009; Robertson {\\it et al} 2009; Springel 2009). The adoption of artificial viscosity by SPH and the limited spatial resolution for AMR, were the reasons for the more significant differences between the two approaches; if however both sources of error are tuned appropriately (e.g. by reducing the action of artificial viscosity outside shocks in SPH, or by increasing the mesh resolution in AMR codes), the two approaches produce consistent results.   \\bigskip   High resolution AMR simulations (such as the ENZO simulations presented in this paper) can provide an accurate representation of the cosmic gas dynamics in galaxy clusters, achieving a very broad dynamical range. Recent works showed that opportune refinements (e.g. Iapichino \\& Niemeyer 2008, Vazza et al. 2009; Maier et al.2009; Zhu, Feng \\& Fang 2010) allow for efficiently studying the onset and evolution of chaotic motions in the ICM.  Yet, some interesting information of Lagrangian nature cannot be followed by these methods. One example is the distribution of metals in the ICM due to diffusion and turbulent mixing which cannot be followed, unless the conservation equations are self-consistently implemented for every metal species, with a sizable increase in algorithmic complexity and computational expense. However, it is possible to treat metals in a post-processing stage by following their spatial evolution through that of mass-less particles (tracers), which move along with the baryon gas.  An other interesting case is the advection of cosmic rays (CR) particles in the ICM. Shocks, turbulence, collisions between high energy hadrons and AGN/supernovae activities are expected to inject relativistic particles in the ICM over cosmological time-scales (e.g. Brunetti \\& Lazarian 2007), whose interplay with diffuse $\\sim \\mu G$ ICM magnetic fields is responsible for the diffuse non-thermal radio emissions observed in galaxy clusters (e.g. Ferrari et al.2008 for a recent review). If the energy stored in CR is much smaller than the thermal ICM (as shown by recent observations of the central regions of clusters, e.g. Reimer 2004; Aharonian et al.2008; Brunetti et al.2007; THE MAGIC collaboration et al.2009), the spatial evolution of CR can be studied through that of passive tracers. This applies as long as the effect of CR diffusion is negligible and the evolution of CR particles can be regarded as the simple advection problem of tracers in the evolving ICM.    \\bigskip  The objective of the present work is to investigate the mixing properties of the ICM with an appropriate numerical recipe. For this purpose high algorithmic accuracy and high spatial and mass resolution are needed. These requirements are well provided by ENZO, which is a high-order and cosmological AMR Eulerian code ( e.g. Norman et al. 2007). Here we present ENZO simulations with two important customizations: first, an additional refinement criterion is added to increase the spatial resolution on propagating shocks (Vazza et al.2009); second, since a pure Eulerian method cannot provide the details on the behavior of each specific fluid element (e.g. its trajectory and velocity), we explicitly follow the the advection of a large number of Lagrangian passive (mass-less) {\\it tracer particles}, which move according to the 3--D velocity field of the gas.  The paper is organized as follows: in Sect.\\ref{sec:methods} we present the details of the simulation runs for this project, and the method to inject and follow tracers. In Sect.\\ref{sec:convergence} we discuss the result of several convergence tests to assess the reliability of tracers simulation with different possible setups. In Sect.\\ref{sec:results} we present astrophysical results of tracers, studying in particular the average transport properties of tracers in all simulated clusters by discussing the evolution of the pair dispersion statistics (Sect.\\ref{subsec:dispersion}); the morphology and the evolution of radial mixing in merging and non merging clusters through the evolution of tracers initially located in spherical shells (Sect.\\ref{subsec:mix}); the spreading of metal tracers in the ICM, studying toy models of metal injections from galaxies (Sect.\\ref{subsec:metals}); the simulated emissivity from the broadened Fe XXIII line from metal tracers and its dependence on the the dynamical history of clusters (Sect.\\ref{subsec:iron_line}). Our conclusions are given in Sect.\\ref{sec:conclusions}.  \\bigskip    \\begin{table} \\label{tab:clusters} \\caption{Main characteristics of the simulated clusters. Column 1: identification number. Columns 2: total mass ($M_{\\rm DM}+M_{gas}$) inside the virial radius.  Columns 3: virial radius. Column 4: dynamical state at $z=0$ (qualitative classification); MM=major merger for $z<1$; mm=only minor mergers for $z<1$, rr=visually relaxed at $z \\approx 0$. Cluster H1d is the same as H1, but re-simulated with mesh refinement on gas/DM density only.} \\begin{center} \\begin{tabular}{c|c|c|c|c} ID & $M_{vir}[10^{14} M_{\\odot}/h]$ & $R_{\\rm v}[kpc/h]$ & dynamical state\\\\ \\hline  H1 & 3.10 & 1890 & mm    \\\\  % H1d & 3.10 & 1890 & mm  \\\\ H2 & 3.05 & 1810 & MM   \\\\   % H3 & 2.95 & 1710 & rr   \\\\  % H4 & 2.63 & 1738 & rr    \\\\ % H5 & 2.41 & 1703 & MM    \\\\  % H6 & 2.32 & 1593 & MM    \\\\  % H7 & 2.14 & 1410 & mm   \\\\   %   \\end{tabular} \\label{tab:char} \\end{center} \\end{table}    \\begin{figure*} \\includegraphics[width=0.3\\textwidth]{images/map_multi0001.ps} \\includegraphics[width=0.3\\textwidth]{images/map_multi0100.ps} \\includegraphics[width=0.3\\textwidth]{images/map_multi594.ps} \\caption{An example of the advection of tracers within a forming galaxy cluster. Contours: projected gas density across a volume of side 10 Mpc/h. Blue points: projected positions of $N \\approx 10^{6}$ at $z=21$, $z=3$ and $z=0$.} \\label{fig:map1} \\end{figure*} ",
        "conclusions": "\\label{sec:conclusions}  In this work we studied the mixing properties of the ICM in simulated galaxy clusters, using tracer particles that are advected by the gas flow during the evolution of cosmic structures. This approach allowed us to obtain valuable information on the Lagrangian properties of the ICM, and  allowed progress in our understanding of the transport processes in Eulerian simulations of galaxy clusters. Seven galaxy clusters with total masses in the range $M \\sim 2 - 3 \\cdot 10^{14}M_{\\odot}/h$ and different dynamical states were simulated with an updated version of the ENZO code (Norman et al.2007), where an additional refinement criterion was designed to reach good spatial accuracy at shock waves up to $\\sim 2 R_{vir}$ from the center of galaxy clusters (Vazza et al.2009). In a post-processing, we injected and followed passive tracers, and tracked their spatial evolution with high time resolution. Numerical tests were presented and discussed to find the optimal strategies to inject tracers, update their positions in time and sample the galaxy cluster volume (Sect.\\ref{sec:convergence}). We then applied this tracer technique to some interesting problems concerning the mixing of the ICM.  First, we showed that the pair dispersion statistics has a well constrained behavior in time: for most of the cosmic time the separation between close couples of tracers shows a very regular evolution in all clusters, consistently with the basic $\\propto t^{3/2}$ scaling expected for a simple turbulent model (Sect.\\ref{subsec:dispersion}). This finding is well consistent with the independent measurement of 3-D velocity power spectrum for the same clusters, and confirms that a sizable amount of complex subsonic motions correlated on cluster scales account for 10-30 per cent of the thermal energy inside $R_{vir}$ of our simulated clusters.  Second, we quantified the degree of radial mixing of the ICM by following the evolution of tracers injected in regular shells within the cluster volume at $z \\approx 1$. The radial mixing morphology was found to be strongly dependent on the dynamical history of every cluster. Large plumes (with the size of $\\sim Mpc$) where gas matter is efficiently mixed were found for a major merger cluster. On the other hand more regular and concentrated mixing profiles were measured for more relaxed clusters.  Third, we assumed a simple model of metal injection from galaxies identified in the simulations around the clusters and studied the metal enrichment of the ICM by tracking the advection of ``metal'' tracers (tracers with a different iron content as a function of their galaxy of origin) deposited in the cluster volume. In this case also, the dynamical state of galaxy clusters and its matter accretion history were found to affect the final distribution of metal tracers at evolved epochs. In particular, the simulated clusters showed a hint of dichotomy in metallicity distributions between clusters experiencing only minor mergers and clusters with major mergers, with the latter showing a broader distribution of metals.  Finally, we simulated the emissivity of the Fe XXIII line from the center of our clusters, using the 3--D distribution of metals produced through our metal tracer injection (model 'c'). We reported significant departures from the simplified assumption of a constant or a radially symmetric metallicity profile, particularly for merger systems. This stresses the importance of studying in detail the connection between the metal enrichment history and the matter accretion history of the ICM.  \\bigskip  In conclusions, we believe that the coupling of high resolution Eulerian numerical simulations and a tracer technique represents a very powerful tool to study crucial open problems concerning the mixing properties of the diffuse gas in galaxy clusters, and the injection and the advection of cosmic rays particles in galaxy clusters, which will be subject of future works."
    },
    "0910/0910.5737_arXiv.txt": {
        "abstract": "{} {We study the observational signature of flux emergence in the photosphere using synthetic data from a 3D MHD simulation of the emergence of a twisted flux tube.} {Several stages in the emergence process are considered. At every stage we compute synthetic Stokes spectra of the two iron lines Fe I 6301.5 {\\AA} and Fe I 6302.5 {\\AA} and degrade the data to the spatial and spectral resolution of Hinode's SOT/SP. Then, following observational practice, we apply Milne-Eddington-type inversions to the synthetic spectra in order to retrieve various atmospheric parameters and compare the results with recent Hinode observations.} {During the emergence sequence, the spectral lines sample different parts of the rising flux tube, revealing its twisted structure. The horizontal component of the magnetic field retrieved from the simulations is close to the observed values. The flattening of the flux tube in the photosphere is caused by radiative cooling, which slows down the ascent of the tube to the upper solar atmosphere. Consistent with the observations, the rising magnetized plasma produces a blue shift of the spectral lines during a large part of the emergence sequence.} {} ",
        "introduction": "Observational studies of magnetic flux emergence in the photosphere with the help of Stokes parameters started with the work of \\citet{Brants:1985a, Brants:1985b} and \\citet{Zwaan_etal:1985}. Subsequent studies showed that active regions grow by the successive emergence of magnetic flux fragments, each containing only a small fraction of the total flux of the active region \\citep{Strous_etal:1996, Strous_Zwaan:1999}. Emerging flux regions (EFRs) of various sizes lead to magnetic complexes in a range of scales \\citep{Harvey:1993}, from the largest active regions ~\\citep[up to $4\\times 10^{22}$ Mx;][]{Zwaan:1987} to small ephemeral regions~\\citep[with fluxes down to $10^{19}$ Mx;][]{Hagenaar:1999}.  While the initial studies were restricted to spectropolarimetric observations of Stokes-$I$, Stokes-$V$ and continuum images \\citep{Brants:1985a, Brants:1985b, Strous:1994}, the full Stokes vector has been used by \\citet[][]{Lites:1998} and \\citet[][]{Kubo:2003} for photospheric studies and by \\citet[][]{Solanki:2003nat} and \\citet{Lagg:2004, Lagg:2007} for chromospheric investigations. Recently, observational studies of flux emergence with high spectral and spatial resolution have been carried out \\citep{Centeno:2007, Cheung:2008, Ishikawa:2008, Martinez:2007,Okamoto:2008} with the spectropolarimeter \\citep[SP;][]{Lites:2001} which is part of the focal-plane package of the $50$ cm Solar Optical Telescope \\citep[SOT;][]{Tsuneta:2008} onboard the Hinode spacecraft. The emerging flux regions show horizontally oriented flux elements, which are generally regarded as the tops of emerging loops. The reported field strength of these horizontal flux elements varies somewhat among observational studies : $500\\pm 300$ G \\citep{Brants:1985a, Brants:1985b}; $200 < |B| < 600$ G \\citep{Lites:1998, Sigwarth:2000}; $400 < |B| < 700$ G \\citep{Kubo:2003}; $~ 650$ G \\citep{Okamoto:2008}. The rising flux elements have an upward velocity of $\\leq 1$ km/s. The estimates of the rise velocity are : $0.5$ km/s \\citep{Brants:1985a, Brants:1985b, Strous_etal:1996}; $\\sim 1$ km/s \\citep{Lites:1998}; $< 1$ km/s \\citep{Kubo:2003}; 0.3 km/s \\citep{Okamoto:2008}.   \\citet{Lites:2005} and \\citet{Okamoto:2008} have shown that some flux elements emerging in the photosphere exhibit a twisted field (cf. \\cite{Wiegelmann:2005}, who found newly emerged loops to harbour considerable electrical currents).~\\citet{Okamoto:2008,Okamoto:2009} (hereafter OK2008 and OK2009) reported observations of the emergence of a coherent helical flux tube.  A number of theoretical studies has addressed magnetic flux emergence from the convection zone to the solar atmosphere. While some approaches treat the convection zone and overlying atmosphere as plane-parallel hydrostatic layers \\citep[e.g.][etc]{Shibata:1989, Abbett:2003, Archontis:2004, Magara:2006}, others include the effects of convective flows \\citep[e.g.][]{Cheung:2007,Martinez:2008,Isobe:2008}.  In this paper we present for the first time detailed synthetic observational properties from 3D MHD simulations of an emerging, initially horizontal, twisted flux tube including its interaction with granular convective flows. We perform spectral line synthesis and inversion at spectral and spatial conditions similar to Hinode's SP/SOT. The results are discussed and some emergence properties are compared with the observations of OK2008 and OK2009, even though it is clear that there are limits to such a comparison, arising from the fact that observations and simulations do not occur in the same solar environment. Nonetheless, we believe that the similarities or differences revealed by a comparison will provide fresh insight into both simulations and observations. ",
        "conclusions": "\\label{sec:4}   Figure~\\ref{fig1} shows a time sequence of the magnetic field maps retrieved by the inversion. Clearly, the emerging field expands laterally with time. Similar expansion has been reported in observations of the emergence of a flux rope in an active region by OK2008 and OK2009 (see Fig. 2 in OK2008 during the first half of the rise of the observed flux rope).  The orientation of the horizontal components of the field is indicated by arrows. The twist of the emerging flux tube is revealed by the orientations of the field lines, although the apparent structure is deformed by the influence of the granulation. This orientation changes along the emergence sequence (see the difference between the two snapshots at $t=11.3$ min and $t=16.1$ min).   At $t=11.3$ min the spectral lines sample the upper part of the emerging flux tube, where the field lines are predominantly oriented from upper left to lower right in the plotted frame. Later, at $t=16.1$ min, the flux tube has emerged further and the spectral lines sample its lower part, where the field lines in this twisted structure are oriented from the lower left to the upper right. A change in the field orientation has been reported in the observed event by OK2008 and OK2009 although the observed event takes place in an active region.   The mean value of the horizontal component of the magnetic field in the region outlined with the blue contour (near the polarity inversion line) in the upper left panel of Fig.~\\ref{fig1} is $550$ G at $t=11.3$ min. This value is close to the one ($650$ G) reported by OK2008 and OK2009. The horizontal field strength exhibits higher values in the close vicinity of the polarity inversion line and decreases with time along the emergence sequence shown in Fig.~\\ref{fig1}. The upward velocity of the emerging flux rope in the same region outlined in blue is $670$ m/s. This value is larger than the $300$ m/s reported in OK2008 and OK2009 and indicates a faster emergence rate of the simulated flux rope. The mean value of the magnetic field inclination with respect to the vertical in the region outlined in blue is nearly zero degrees (zero corresponds to a horizontal field) with a standard deviation of about 30 degrees. In OK2009, it is reported that the mean inclination near the polarity inversion region is nearly zero, with a scattering of the order of $\\pm 10$ degrees. The mean vertical magnetic field in the same outlined region is nearly zero. The one reported in OK2009 is 0 $\\pm 200$ G.   The total unsigned magnetic flux through the maps in Fig.~\\ref{fig1} is calculated from the inversion results as: Flux = $A \\displaystyle\\sum_{i} |B_{z,i}| f_{i}$, where $A=$ pixel area, $B_{z,i}=B_{z}$ in pixel $i$ and $f_{i}=$ filling factor in pixel $i$. For the four snapshots a-d in Fig.~\\ref{fig1} we obtain the following values: $1.15 \\times 10^{20}$ Mx, $1.42 \\times 10^{20}$ Mx, $1.45 \\times 10^{20}$ Mx and $1.05 \\times 10^{20}$ Mx, respectively. The increase of the unsigned flux in the first three snapshots is due to flux emergence. Subsequently, the decrease is due to flux cancellation of opposite-polarity flux elements. These elements are formed by the fragmentation of the tube under the action of flux expulsion \\citep{Voegler:2005}. The encounter of opposite polarity flux elements in the intergranular lanes leads to the decrease of the unsigned flux.    2D vertical cuts through the MHD cubes perpendicular to the axis of the original flux rope are shown in Figs.~\\ref{fig2} and~\\ref{fig3} at t=11.3 min and 16.1 min, respectively. They lie along the horizontal coordinates represented by the black line in the upper left panel of Fig.~\\ref{fig1}.   The rising flux tube flattens considerably after reaching the solar surface. The upper parts of the flux tube (above the solar surface) rapidly cool down radiatively. Therefore, the plasma becomes denser and loses its buoyancy. On the other hand the bulk of the flux tube keeps rising. This is due to the fact that the rising sub-surface parts of the flux tube keep pushing the plasma upward above the surface \\citep{Cheung:2007}. The material above the surface also expands laterally due to the strong stratification of the photospheric layers and partly descends back into the convection zone at the sides of the flux tube. We see a blue shift during most of the emergence sequence at the height where Fe I 6301.5 {\\AA} and Fe I 6302.5 {\\AA} are formed, except for downflow lanes along the flux rope where we find a red shift \\citep[][]{Cheung:2007,Yelles:2008}. A blue shift has been observed as well by OK2008 during a large part of the emergence sequence.       Figure~\\ref{fig2} shows that the flux rope has a relatively regularly twisted core with a field strength approaching 1 kG near and below the solar surface. Note that the Stokes-$V$ signal is formed over a height-range sampling mainly sub-kG field. Convective flows affect the outer parts of the emerging flux tube more strongly, so that the field configuration becomes more complex than in the central part. This is best seen in the field inclination (lowest panel of Fig.~\\ref{fig2}) and is mainly restricted to the sub-surface layers. At the height of the line formation the field is rather homogeneous and the opposite polarities are separated (see also the top left panel in Fig.~\\ref{fig1}).      \\begin{figure} % \\includegraphics[width=0.9\\textwidth,bb= 75 123 760 590]{vercut_emergence22_5.ps} \\caption{Vertical cut through the simulation box at the location shown by the black line in the upper left panel in Figure~\\ref{fig1} at $t= 11.3$ min. The panels represent, from top to bottom, the temperature fluctuations with respect to the horizontal average, vertical velocity (blue represents upflow), magnetic field strength and magnetic field inclination. The solid black lines indicate the location of $\\tau_{5000}=1$. The dash-dotted lines show the height at which the contribution function of Stokes-$I$ reaches its maximum, and the dashed lines indicate the same for Stokes-$V$.} \\label{fig2} \\end{figure}    \\onlfig{3}{ \\begin{figure} % \\includegraphics[width=0.9\\textwidth,bb= 75 123 760 590]{vercut_emergence58_5.ps} \\caption{Similar to Figure~\\ref{fig2} at $t= 16.1 $ min.} \\label{fig3} \\end{figure} }    At $t=16.1$ min (Fig.~\\ref{fig3}), one can still see the downflows beside the flux tube, where the field is predominantly vertical. The flux tube has become wider, and there are patches of downflowing plasma above the solar surface in the body of the flux tube (e.g. near the abscissa $x=5500$ km). The horizontally oriented core of the flux tube has greatly shrunk. Part of the flux has been advected to downflow lanes and concentrated to kilo-Gauss field strength. This produces a strong circularly polarized signal correlated with downflow lanes in the emergence area. At this point in time the flux rope is starting to loose its coherence even at the height of the line formation (Fig.~\\ref{fig3}), with multiple islands of both polarities being present.  By the time of the last plotted snapshot (Fig.~\\ref{fig1}d) many islands of opposite polarities are visible. Nonetheless, a predominance of positive polarities in the upper and negative polarities in the lower part of the frame can be discerned. In vertical cuts similar to Figs.~\\ref{fig2} and~\\ref{fig3} we show how the inhomogeneity of the magnetic structure slowly propagates upwards."
    },
    "0910/0910.0865_arXiv.txt": {
        "abstract": "{ Neptune Trojans and Plutinos are two sub-populations of Trans-Neptunian Objects located in the 1:1 and the 3:2 mean motion resonances with Neptune, respectively, and therefore protected form close encounters with the planet. However, the orbits of these two kinds of objects may cross very often, allowing a higher collisional rate between them than with other kinds of Trans-Neptunian Objects and a consequent and size distribution alteration of the two sub-populations.  Observational colors and absolute magnitudes of Neptune Trojans and Plutinos show that: i) there are no intrinsically bright (large) Plutinos at small inclinations; ii) there is an apparent excess of blue and intrinsically faint (small) Plutinos; and iii) Neptune Trojans possess the same blue colors as Plutinos within the same (estimated) size range do.  For the present sub-populations we analyze the most favorable conditions for close encounters / collisions to occur and address if there could be a link between those encounters and the sizes and/or colors of Plutinos and Neptune Trojans. We also perform a simultaneous numerical simulation of the outer Solar System over 1~Gyr for all these bodies in order to estimate their collisional rate.  We conclude that orbital overlap between Neptune Trojans and Plutinos is favored for Plutinos with high libration amplitudes, high eccentricities and low inclinations. Additionally, with the assumption that the collisions can be disruptive creating smaller objects not necessarily with similar colors, the present high concentration of small Plutinos at low inclinations can thus be a consequence of a collisional interaction with Neptune Trojans and such hypothesis should be further analyzed. } ",
        "introduction": "Trans-Neptunian Objects (TNOs), also known as Kuiper Belt Objects (KBOs), are a population of small and primitive icy bodies orbiting (mostly) beyond Neptune. Their study is one of the most significant ways to obtain information on the early ages of the Solar System.  Based on some distinct dynamical properties, the TNOs can be subdivided in several different sub-populations, often also called families or groups. Subdivide them as a function of  their physical properties seems to be far more complex. TNOs have surface colors so diverse that can go from blue/neutral ({\\it i.e.} solar-like) to extremely red. A possible explanation for the wide variety of colors was originally proposed by \\cite{1996AJ....112.2310L} with the collisional resurfacing model. In this model, the competition between surface reddening, due to cosmic-ray bombardment, and a resurfacing with frozen material (assumed to be bluer) withdrawn from beneath its crust by impact collisions could be responsible for the observed wide range of surface colors. As the model seems to fail, at least in its simple form, more complex forms have been proposed. Namely, collisional resurfacing combined with cometary activity \\citep{Delsanti+04}, and collisional resurfacing with layered reddening and even re-bluing by cosmic-rays \\citep{2002P&SS...50...57G}. An alternate idea proposes that surface colors are primordial \\citep[][and references therein]{TRC03}. More recently, \\cite{2009Icar..199..560G} shows that an object could lose its redness simply by ice sublimation (without resurfacing). Such could explain the lack of red colors among TNOs that evolved into Jupiter family comets due to their closeness to the Sun. Yet, as the aforementioned work acknowledges, the existence of the blue/neutral TNOs at heliocentric distances where ices do not sublimate suggests the existence of other coloring mechanisms. Our understanding on the origin and eventual alteration of TNOs colors is still very limited and, up to the present, none of these approaches lead to a fully consistent explanation for the color diversity \\cite[for a review see][]{2008ssbn.book...91D}.  Among the TNOs, two sub-populations caught our attention: the Neptune Trojans and the Plutinos. Neptune Trojans are the small bodies trapped in a 1:1 mean motion orbital resonance with Neptune, also with orbital eccentricities lower than 0.1. Roughly they co-orbit with Neptune concentrated $60^{\\circ}$ ahead and $60^{\\circ}$ behind Neptune's position. \\cite{2005ApJ...628..520C} proposed that Neptune Trojans formed by in-situ accretion from small-sized debris once Neptune's migration has stopped. On the other hand, other works indicate they can survive planetary migration \\citep{2002Icar..160..271N, 2004Icar..167..347K} and, more recently, due to the discovery of a large population of high inclination Neptune Trojans, \\cite{2009AJ....137.5003N} sustain they were captured during planetary migration Compared to TNOs Neptune Trojans seem to be quite small in size (diameters $D<100$ km) and also slightly blue.  At the time of our analysis 6 Neptune Trojans were known\\footnote{See http://cfa-www.harvard.edu/iau/lists/NeptuneTrojans.html}, all of them librating around the Lagrangian point $L_4$. %  Plutinos are those trapped in a 3:2 mean motion orbital resonance with Neptune, with eccentricity values ranging between 0.1 and 0.3. Unlike Trojans, the colors of Plutinos vary from blue/neutral to the very red, and their sizes range from a few tens of~km to a few thousands (note that Pluto is a Plutino). Roughly, Plutinos possess semi-major axes within $39 < a < 40.5$ AU. Through long-term dynamical evolution studies \\cite{2007Icar..189..213L} identified 98 Plutinos and that will be our reference.  Being locked at the 3:2 resonance, Plutinos can periodically cross the orbit of Neptune without colliding with it. However, this protection from collisions is not possible for Neptune Trojans and Plutinos. A first look at the geometry of these two populations suggests that they might even collide frequently.  The analysis of the collisional resurfacing model by \\cite{2003Icar..162...27T} and \\cite{2003EM&P...92..233T} found that Plutinos were significantly more affected by collisions than other TNOs. Since, observationally, Plutinos did not exhibit bluer colors than all other TNOs taken as a whole, the aforementioned work strongly argued against the collisional resurfacing model, at least as major cause for the color diversity of TNOs. Neptune Trojans were not included in \\citeauthor{2003Icar..162...27T}'s simulations, though, since they were not known at the time. Notwithstanding the arguments against it, the collisional resurfacing that was analyzed was a very simple model. For instance, the possibility of collisional disruption of differentiated bodies producing fragments or ruble-piles with surface colors distinct from those of the parent bodies has not been taken into account in the collisional resurfacing models yet. Naturally, the possible outcomes of this scenario would be far more complex and our understanding of the physics of collisional processes between the icy TNOs is still limited \\citep[see review by][]{2008ssbn.book..195L}. Nonetheless, both a collisional family of TNOs associated with (136108) Haumea and a collisional system of satellites associated with Pluto have been detected \\citep{2007Natur.446..294B, 2006Natur.439..946S, 2005Sci...307..546C}.   Surveys show a power-law for the size distribution of TNOs with slope change at $D\\sim100$~km  \\citep{2004AJ....128.1364B}, which is most consistent with TNOs being gravity dominated bodies with negligible material strength. Objects larger than $\\sim 100$~km are difficult to disrupt, hence likely primordial, and objects smaller than this are expected to be shattered due to disruptive collisional evolution \\citep{2008ssbn.book..293K, 2005Icar..173..342P}.  \\cite{Elia+08} studied the collisional evolution of Plutinos considering only Plutino-Plutino collisions. They infer that any eventual slope change at  $D\\sim40-80$~km on the power-law size distribution of Plutinos should be primordial and not of collisional origin. Yet, they find that collisional populations may form from the breakup of objects larger than $100$~km, and if those populations form at low inclinations their fragments will likely stay in the resonance.  On Sect. \\ref{obs} we will analyze the observational data relative to Neptune Trojans and Plutinos. We will discuss two particular observational properties: 1) there is an apparent ``excess'' of small blue Plutinos and these happen to be also in the same size range and to possess the same colors as the known Neptune Trojans, and 2) there is a concentration of small Plutinos, and absence of large ones, at low orbital inclinations. The eventual collisional interaction between Plutinos and Neptune Trojans has not been analyzed in previous works. Since their geometry points to the existence of mutual collisions, the existence of an eventual link between the colors and sizes of Neptune Trojans and Plutinos and their interactions/collisions motivated us to this work. On Sect. \\ref{dynamiceq} we will study the dynamics of Neptune Trojans and Plutinos, separately, focusing on their orbits around the Lagrangian point $L_4$. On Sect. \\ref{numsim} we will discuss the possibility of collisions between Neptune Trojans and Plutinos and the best conditions for this to happen. Finally, on Sect. \\ref{conclusions} we will present our conclusions. ",
        "conclusions": "\\label{conclusions}  In this work we aimed to verify if both Plutino and Neptune Trojan populations can be merged, and how frequently does that occur, and examine if there could be a connection between such possibility and the observed properties of the two (sub-) populations of TNOs in question. We analyzed the available colors and absolute magnitudes of Plutinos and Neptune Trojans. The nonexistence of significant albedo diversity among the large majority of these objects is assumed, hence we interpreted absolute magnitudes as an equivalent to to object size. We find that:  \\begin{description} \\item[(i)] there are no intrinsically bright (large) Plutinos at small inclinations; \\item[(ii)] there is an apparent excess of blue and intrinsically faint (small) Plutinos; \\item[(iii)] Neptune Trojans possess the same blue colors as Plutinos within the same (estimated) size range do. \\end{description}  From these results, we have hypothesized that there might have been some strong collisional interaction between Neptune Trojans and (1) the presently small Plutinos --- with the assumption that such collisions created only blue colored objects ---, or (2) the low inclined Plutinos --- with the assumption that such collisions created both blue and red objects. Paradoxically, the previous two opposite assumptions both seem to have observation support among other TNOs.  In order to differentiate between these two scenarios we performed a numerical simulation over 1~Gyr of the future evolution of the outer Solar System composed by 5 planets, 6 Neptune Trojans, and 98 Plutinos.  By plotting simultaneously the orbits of the Neptune Trojans and Plutinos in a co-rotating frame with Neptune, it becomes clear that there is a large overlap of the orbits of the two kinds of TNOs. A more detailed analysis revealed, however, that close encounters with the Neptune Trojans are favored for Plutinos with high libration amplitudes, high eccentricity values, and low inclined orbits. Tables \\ref{tab:tab_3} and \\ref{tab:tab_4} show the libration amplitudes and periods computed for these objects.  After 1~Gyr of numerical simulations we registered 2 Trojan-Plutino, 2 Trojan-Trojan, and 10 Plutino-Plutino close approaches, i.e. less than $\\sim$ 300\\,000 km. Since we have used the currently known objects, Trojans are much less numerous than Plutinos in our simulation. If the number of objects among each population is similar then we roughly infer that Trojan-Plutino and Trojan-Trojan ``collisions'' can be about three and fifty times more frequent, respectively, than Plutino-Plutino collisions.  The collision rates between Neptune Trojans and between Neptune Trojans and Plutinos might explain both why we observe a small number of Neptune Trojans and why there is an absence of large Neptune Trojans: a strong collisional evolution possibly played an important role by shattering and/or depleting Neptune Trojans.  Our results also show that Plutinos in low-inclined orbits have more chances of colliding with Neptune Trojans. This result gives no support for a (mutual) collisional origin of the equal-sized and equal-colored Neptune Trojans and small Plutinos, as the latter are equally spread in inclination (see Sect. \\ref{obs}, scenario~\\#1). On the other hand, it gives plausibility to the origin of the concentration of small Plutinos with $i<8-13^\\circ$ as being a consequence of some collisional interaction with Neptune Trojans(see Sect. \\ref{obs}, scenario~\\#2).  During our numerical simulations we also observed that the orbit of Trojan 2001QR322 became unstable as well as the orbits of 10 more Plutinos. The stability of Neptune Trojans appears to be enhanced for high inclination values \\citep{2007MNRAS.382.1324D}, and also for small libration amplitudes. On the other hand, Plutinos seem to be essentially pushed out of their resonance by Pluto, in conformity with the results of \\citet{1999AJ....118.1873Y} and \\citet{Nesvorny_etal_2000}.  This work's results were derived for objects considered as test particles (except for Pluto) and the number of ``collisions'' were extrapolated from the amount of close encounters after each integration step-size. Accurate estimations for the amount of collisions between Neptune Trojans and Plutinos can only be made with the inclusion of mutual gravitational interactions between all objects, and using the real spatial and size distributions of these objects.  Nonetheless, our sketchy analysis indicates that under certain assumptions, which have parallel in what has been inferred for other TNOs, shattering collisions involving Neptune Trojans might have played a crucial role on the creation of the size-inclination asymmetries observed among Plutinos. We recall that we have disregarded the possibility that these asymmetries might have been created by collisional interaction with some other sub-population of TNOs than the Neptune Trojans in view of the results obtained by \\cite{2003Icar..162...27T} and \\cite{2003EM&P...92..233T}. If those works come to be revealed as inadequate approximations of the relative collision rates suffered by TNOs our inference on a possible cause for the size-inclination distribution of Plutinos looses its ground. Further, if Neptune Trojans and low inclined Plutinos do not possess identical spin rate {\\it versus} size distributions, which should be distinct from the higher inclined Plutinos, then our suggestions cannot hold either \\citep[e.g.][]{1981A&A...104..159F}. More detailed studies on the interaction between Neptune Trojans and Plutinos should be attempted."
    },
    "0910/0910.5392_arXiv.txt": {
        "abstract": "We develop a new model of the fluctuation dynamo in which the magnetic field is confined to thin flux ropes advected by a multi-scale flow which models turbulence. Magnetic dissipation occurs only via reconnections of flux ropes. The model is particularly suitable for rarefied plasma, such as the Solar corona or galactic halos. We investigate the kinetic energy release into heat, mediated by dynamo action, both in our model and by solving the induction equation with the same flow. We find that the flux rope dynamo is more than an order of magnitude more efficient at converting mechanical energy into heat. The probability density of the magnetic energy released during reconnections has a power-law form with the slope $-3$, consistent with the Solar corona heating by nanoflares. We also present a nonlinear extension of the model. This shows that a plausible saturation mechanism of the fluctuation dynamo is the suppression of turbulent magnetic diffusivity, due to suppression of random stretching at the location of the flux ropes. We confirm that the probability distribution function of the magnetic line curvature has a power-law form suggested by \\citet{Sheck:2002b}. We argue, however, using our results that this does not imply a persistent folded structure of magnetic field, at least in the nonlinear stage. ",
        "introduction": "Dynamo action, i.e., the amplification of magnetic field by the motion of an electrically conducting fluid (plasma), is the most likely explanation for the omnipresence of astrophysical magnetic fields. Ohmic dissipation, however small, is essential in order to achieve the development of the dynamo eigensolutions and to smooth out the spatial variations of the magnetic field. The evolution of the magnetic field $\\mathbf{B}$ embedded in a velocity field $\\mathbf{u}$ is governed by the following closed equation: \\begin{equation}\\label{induction} \\pderiv{\\mathbf{B}}{t}=\\nabla \\times (\\mathbf{u} \\times \\mathbf{B}) +\\widehat{\\cal L}\\mathbf{B}, \\end{equation} where $\\widehat{\\cal L}$ is an operator describing the magnetic dissipation.  In rarefied astrophysical plasmas, such as the Solar co\\-ro\\-na, hot gas in spiral and elliptical galaxies, galactic and accretion disc halos, and laboratory plasmas, an important (or even dominant) mechanism for the dissipation of magnetic field is the reconnection of magnetic lines rather than magnetic diffusion \\citep{priest:2000}, the latter modelled with $\\widehat{\\cal L}=\\eta\\nabla^2$ (if $\\eta=\\mbox{const}$). Discussions of astrophysical dynamos often refer to magnetic reconnection, but attempts to include features specific of magnetic reconnection into dynamo models are very rare \\citep[see however ][]{Blackman:1996}. On the other hand, theories of magnetic reconnection (and the resulting estimates of the plasma heating rate) rarely, if ever, refer to the dynamo action as the widespread mechanism maintaining magnetic fields. This paper attempts to bridge the gap between the two major topics of astrophysical magnetohydrodynamics (dynamos and reconnections) by developing a dynamo model which explicitly incorporates magnetic reconnections.  \\begin{figure}[h!] \\begin{center} \\psfrag{x}[c]{$k$} \\psfrag{y}[c][l]{$\\widehat{\\cal L}_k$} \\psfrag{t}[l]{${2\\pi}/d_0$} \\psfrag{u}[c]{$k^2$} \\psfrag{v}[c]{$k^4$} \\includegraphics[width=0.4\\textwidth]{fourie_scheme.eps} \\caption{\\label{fourier}  A schematic representation of the magnetic dissipation operator $\\widehat{\\cal L}$ in Fourier space: usual diffusion $\\widehat{\\cal L}_k \\propto k^2$ (dash-dotted), hyperdiffusion $\\widehat{\\cal L}_k \\propto k^4$ (dashes) and reconnections at a scale $d_0$ as described by our model (solid).} \\end{center} \\end{figure}  The nature of the dissipation mechanism is important for the dynamo action as it affects the growth time of magnetic field, its spatial form and the rate of plasma heating by the electric currents. For example, dynamo action with hyperdiffusion, $\\widehat{\\cal L}=-\\eta_1\\nabla^4$ (and with a helical $\\mathbf{u}$) has larger growth rate and stronger steady-state magnetic fields than a similar dynamo based on normal diffusion \\citep{Brandenburg:2002}. This is not surprising, as the hyperdiffusion operator, having the Fourier dependence of $k^4$, rather than the $k^2$ dependence of normal diffusion, has weaker magnetic dissipation at larger scales as shown in Fig.~\\ref{fourier}. This allows the magnetic field to grow unimpeded by dissipation as magnetic dissipation is confined to relatively small regions. The release of magnetic energy in smaller regions (and larger current densities) in hyperdiffusive dynamos may also lead to a higher rate of conversion of kinetic energy to heat via magnetic energy. One of the aims of this paper is to demonstrate that this statement is especially true in the case of magnetic reconnections.  Magnetic hyperdiffusion also appears in the context of continuous models of self-organised criticality in application to the heating of the Solar corona \\citep{SOC}. The aim of such models is to reproduce the observed frequency distribution of various flare energy diagnostics. As we show here, our model exhibits a power-law probability distribution of the magnetic energy release similar to that observed in the Solar corona.  Magnetic reconnection may correspond to an even more extreme form of the dissipation operator than the hyperdiffusion: here magnetic flux tubes dissipate their energy only when in close contact with each other, so that the Fourier transform of $\\widehat{\\cal L}$ should be negligible at all scales exceeding a certain reconnection length $d_0$ (see Fig.~\\ref{fourier}). It is then natural to expect that dynamos based on reconnections (as opposed to those involving magnetic diffusion) will exhibit faster growth of magnetic field, more intermittent spatial distribution and stronger plasma heating. In this paper we consider dynamo action based on direct modelling of magnetic reconnections. For this purpose, we follow the evolution of individual closed magnetic loops in various flows (known to support dynamo action) and reconnect them directly whenever they come into a sufficiently close contact, with appropriate magnetic field directions. First results of our simulations can be found in \\citep{BBSS09}.  \\begin{figure} \\begin{center} \\includegraphics[width=0.35\\textwidth]{insert_scheme.eps} \\caption{\\label{insert_scheme} (Colour online) The algorithm for inserting new trace particles in a stretched (left to right) or contracting (right to left) magnetic flux tube. If the distance between any two trace particles (shown with red/open circles) exceeds a length scale $d$, a new particle is inserted between the two particles shown with a blue/filled circle. The label at each particle represents magnetic field strength at that location } \\end{center} \\end{figure} ",
        "conclusions": "To summarise, we have confirmed that the dynamo action is sensitive to the nature of magnetic dissipation and demonstrated that magnetic reconnections (as opposed to magnetic diffusion) can significantly enhance the dynamo action. We have explored the kinematic stage of the fluctuation dynamo in a chaotic flow that models hydrodynamic turbulence and in the ABC flow, with the only magnetic dissipation mechanism being the reconnection of magnetic lines implemented in a direct manner. In our model, where magnetic dissipation is suppressed at all scales exceeding a certain scale $d_0$, the growth rate of magnetic field exceeds that of the magnetic diffusion-based fluctuation dynamo with the same velocity field. Even when the velocity field of the reconnection-based dynamo is reduced in magnitude as to achieve similar growth rates of magnetic energy density, the rate of conversion of magnetic energy into heat in the reconnection dynamo is an order of magnitude larger than in the corresponding diffusion-based dynamo. Thus, reconnections more efficiently convert the kinetic energy of the plasma flow into heat, in our case with the mediation of the dynamo action. This result, here obtained for a kinematic dynamo, can have serious implications for the heating of rarefied, hot plasmas where magnetic reconnections dominate over magnetic diffusion (such as the corona of the Sun and star, galaxies and accretion discs).  It is intriguing that reconnections play the same role \\citep{Leadbeater:2001,Barenghi:2008} of converting kinetic energy into heat in superfluids and Bose-Einstein condensates, fluids near absolute zero at the opposite end of the temperature spectrum.  Our model can be viewed as a numerical implementation of the elusive limiting regime of infinitely large magnetic Reynolds number, where magnetic dissipation can be safely neglected at all large scales but plays a crucial role at a certain very small scale (we are grateful to Alex Schekochihin for suggesting this idea).  In contrast to the fluctuation dynamo based on magnetic diffusion, the probability distribution function of the energy released in the flux rope dynamo has a power law form not dissimilar to that observed for the Solar flares. This is also true for the nonlinear states of the dynamo.  The reconnection-based dynamo model suggested here can be generalised to include the modification of the velocity field by the Lorentz force. More precisely, magnetic pressure is assumed to be balanced by the gas pressure, so that only magnetic tension needs to be explicitly included into the Navier--Stokes equation. Magnetic tension can readily be calculated in our model where magnetic field is defined only at discrete positions of closed magnetic loops. We suggest two approximations for the Navier--Stokes equation, one designed to model Alfv\\'en waves and the other suitable for the studies of nonlinear dynamos. The former model can be useful in the studies of nonlinear interaction of Alfv\\'en waves and Alfv\\'enic turbulence.  Unlike most -- if not all -- other simulations of the fluctuation dynamo, our computations start with a spatially localised initial magnetic field. This has allowed us to observe that the magnetised region spreads during the kinematic dynamo stage but its size stops growing in the nonlinear stage. This can be naturally interpreted as the suppression of the turbulent magnetic diffusion in the saturated dynamo state. This is broadly equivalent to the reduction of the effective magnetic Reynolds number down to its marginal value (with respect to the dynamo action).  Our model of magnetic field evolution, based on tracing closed magnetic loops can be fruitfully applied in other numerical approaches to magnetohydrodynamics. One of well-known difficulties in the generalisation of smoothed-particle hydrodynamics to include magnetic fields is the implementation of the solenoidality of magnetic field. Quite notably, our approach satisfies the magnetic solenoidality condition perfectly since the modelled magnetic lines are closed at all times. A similar approach may be fruitful in smoothed-particle magnetohydrodynamics codes."
    },
    "0910/0910.1358_arXiv.txt": {
        "abstract": "We present a high-resolution spectrum of a microlensed G dwarf in the Galactic bulge with spectroscopic temperature $T_\\mathrm{eff} = 5600 \\pm 180$ K. This $I \\sim 21$ mag star was magnified by a factor ranging from 1160 to 1300 at the time of observation. Its high metallicity ($\\mathrm{[Fe/H]} = 0.33 \\pm 0.15$) places this star at the upper end of the bulge giant metallicity distribution. Using a K-S test, we find a 1.6\\% probability that the published microlensed bulge dwarfs share an underlying distribution with bulge giants, properly accounting for a radial bulge metallicity gradient. We obtain abundance measurements for 15 elements and perform a rigorous error analysis that includes covariances between parameters. This star, like bulge giants with the same metallicity, shows no alpha enhancement. It confirms the chemical abundance trends observed in previously analyzed bulge dwarfs. At supersolar metallicities, we observe a discrepancy between bulge giant and bulge dwarf Na abundances. ",
        "introduction": "Baade's discovery (1946, 1958) that the central regions of the Galaxy, like other Sb spirals, contained population II stars sparked intense interest in the Galactic bulge. Because it is the closest galactic spheroid to the Sun, we can study the Galactic bulge in unique detail and use it to better understand galaxy formation and evolution.  Proposed senarios for building bulges include mergers and secular dynamical evolution (e.g.\\ \\citealt{kormendy:04}). Discriminating between bulge formation theories hinges on age determinations of bulge stars, a process complicated by crowding, contamination by foreground stars, reddening, and dispersion in both metallicity and distance along the line of sight. Comparing main-sequence photometry of bulge globular clusters with halo globular clusters demonstrated that they are coeval \\citep{ortolani:95}. Examining the color-magnitude diagrams of both bulge fields and clusters shows that the bulge clusters are representative of the general bulge population, the majority of which are old stars ($\\sim10$ Gyr) (e.g.\\ Ortolani et al.\\ 1995; Zoccali et al.\\ 2003). Although \\citet{feltzing:00} agree that the bulge is mostly old, their interpretation of the photometry permits a young metal-rich population of stars.  Evidence for a short star formation phase comes from the chemical composition of bulge stars. The $\\alpha$-elements (O, Mg, Si, Ca, and Ti) are created in the evolution and explosion of massive stars as core collapse supernovae (SNcc) and almost immediately recycled in the interstellar medium (ISM) (e.g.\\ \\citealt{woosley:95}). Because of their longer progenitor lifespan, Type Ia supernovae (SNIa) do not enrich the ISM with substantial amounts of iron until 1-1.5 Gyr after star formation begins \\citep{matteucci:09}. The delayed SNIa iron contribution dilutes the SNcc-enhanced $\\alpha$/Fe ratio until it reaches solar levels \\citep{tinsley:79}. Since bulge giants maintain an elevated $\\alpha$/Fe ratio to high metallicity, the bulge must have evolved faster than the solar neighborhood and undergone an intense burst of star formation early in the history of the galaxy (e.g.\\ \\citealt{matteucci:90,mcwilliam:94}).  Within this framework, the $\\alpha$-elements exhibit different behavior. The explosive $\\alpha$-elements Si, Ca, and Ti appear to track each other \\citep{fulbright:07}. Surprisingly, the hydrostatic $\\alpha$-elements O and Mg do not, even though both are made in SNcc progenitors and released in SN explosions. In particular, Mg remains enhanced in the bulge to supersolar metallicities \\citep{mcwilliam:94,fulbright:07}. The bulge O abundance closely tracks the abundance trends in disk stars and begins declining around [Fe/H]$\\approx-0.3$ dex (e.g.\\ \\citealt{melendez:08}). \\citet{mcwilliam:08} reconciles these trends by invoking Wolf-Rayet winds in metal-rich, massive stars to produce lower mass SNcc progenitors, which \\citet{woosley:95} predict yield less oxygen.  Neutron-capture elements also have the potential to act as clocks. AGB stars are associated with the production of $s$-process elements, while supernovae may be responsible for $r$-process elements. Little information about neutron-capture elements is available in the bulge since the absorption lines are concentrated in the crowded blue portion of the spectrum. Nevertheless, \\citet{mcwilliam:04} observed an enhanced Eu abundance and tentatively suggest a behavior similar to Ca. \\citet{mcwilliam:94} find a subsolar $s$-process abundance, reminiscent of old, metal-poor halo stars, indicating that the $r$-process dominates in the bulge.  Light elements, like Na, contain important information about the chemical evolution of the bulge. Na production is sensitive to the neutron excess \\citep{arnett:71} and its yield is therefore dependent on the metallicity of its SNcc progenitor. Although typically interpreted as SNcc ejectum, Na can also be synthesized in envelope burning of AGB stars. \\citet{lecureur:07} report a flat Na abundance trend in the bulge below solar metallicity, after which it increases abruptly with a dispersion at high-metallicities exceeding measurement error. Lecureur et al.\\ determine that the elevated Na abundance cannot be explained by the mild internal mixing seen in their sample of giants and therefore must reflect the composition of the ISM at its formation.  Our knowledge of the bulge's chemical evolution comes from spectroscopic studies of K and M giants because they are typically the only stars accessible to high-resolution observations due to the far distance and the high degree of interstellar extinction toward the bulge. However, large surveys designed to identify and follow up microlensing events, such as the Optical Gravitional Lens Experiment\\footnote{http://ogle.astrouw.edu.pl/ogle3/ews/ews.html} (OGLE) and the Microlensing Observations in Astrophysics (MOA) collaboration\\footnote{www.massey.ac.nz/$\\sim$iabond/alert/alert.html}, have made possible the high-resolution spectroscopic study of otherwise unobservable bulge dwarfs. \\citet{lennon:96} performed the first spectroscopic abundance determination for a microlensed bulge dwarf, 96-BLG-3, detected by the MACHO collaboration. Based on low-resolution data, they estimated a metallicity of between 0.3 and 0.6 dex. \\citet{minniti:98} obtained a high-resolution spectrum of a dwarf using Keck. \\citet{cavallo:03} presented a preliminary abundance analysis of six microlensed bulge stars, including the Minniti et al.\\ star and two other dwarfs.  There are seven published abundances analyses of microlensed bulge dwarfs based on high-resolution data: OGLE-2006-BLG-265S \\citep{johnson:07}, OGLE-2007-BLG-349S \\citep{cohen:08}, MOA-2006-BLG-099S \\citep{johnson:08}, OGLE-2008-BLG-209S \\citep{bensby209:09}, MOA-2008-BLG-310S and MOA-2008-BLG-311S \\citep{cohen:09}, and OGLE-2009-BLG-076S \\citep{bensby76:09}. Although OGLE-2008-BLG-209S is actually  a sub-giant, we will refer to these seven stars collectively as bulge dwarfs. Additionally, \\citet{bensbyIAU:09} announced that high-resolution spectra have been obtained for six more bulge dwarfs, including a reanalysis of the Cavallo et al.\\ dwarfs, whose detailed abundances will appear in Bensby et al. (in prep.) and Johnson et al. (in prep.).  Studying dwarfs offers several advantages. Dwarfs are subject to different systematic effects than the giants. Mixing on the red giant branch (RGB) can destroy some elements and dredge up others, altering a star's surface composition. Some elements that are difficult to see in giants, like S and Zn, can be measured in dwarfs since their hotter temperatures strengthen these atomic lines while weakening CN. Additionally, we can compare the metallicity distribution function (MDF) for dwarfs and giants.  The first-observed high-magnification bulge dwarf with a high-resolution spectrum had a surprisingly high metallicity of 0.56 dex \\citep{johnson:07}. Subsequent observations revealed more metal-rich dwarfs indicating that high metallicities are common in bulge dwarfs. By contrast, the bulge giant MDF peaks just below solar metallicity (e.g.\\ \\citealt{fulbright:07,cunha:06,rich:05,rich:07,lecureur:07,zoccali:03,zoccali:08}). \\citet{cohen:08} suggest that high-metallicity stars experience a sufficiently high mass-loss rate to cause them to peel off the RGB before undergoing a He flash and become He white dwarfs (WDs). \\citet{zoccali:08} counter that this scenario would produce a decline in the RGB luminosity function, which is not observed in the \\citet{zoccali:03} data. OGLE-2008-BLG-209S, the first bulge sub-giant with a high-resolution spectrum, was found to have a subsolar metallicity, consistent with the giants also located in Baade's window \\citep{bensby209:09}. The detection of metal-poor dwarf OGLE-2009-BLG-076S opposes the idea of a strong selection effect against metal-poor microlensed stars.  Here we present both the high-magnification bulge dwarf MDF and a detailed chemical abundance analysis of the eighth bulge dwarf, OGLE-2007-BLG-514S. ",
        "conclusions": "We performed a detailed chemical abundance analysis of 15 elements in OGLE-2007-BLG-514S, a G dwarf star located in the bulge. Gravitational microlensing magnified this $I\\sim21$ mag star by a factor $\\gtrsim$1000, affording us a rare glimpse of the composition of such a faint star. This brings the number of dwarfs in the bulge with high-resolution spectra and published abundance ratios to eight.  OGLE-2007-BLG-514S boasts a high metallicity with [Fe/H]$=0.33\\pm0.15$ dex, swelling the number of metal-rich bulge dwarfs to three-quarters of the sample. The prevalence of bulge dwarfs at high [Fe/H] conflicts with the wide distribution of bulge giant metallicities centered slightly below solar metallicity. We create a simple model for how a bulge metallicity gradient would affect the bulge giant MDF if the giants had the same spatial distribution as the dwarfs. A K-S test yields a 1.6\\% probability that the bulge dwarf MDF is identical to our model giant MDF. We ascribe this incongruity to one of three possibilities. One option is that metal-rich dwarfs are not becoming metal-rich giants. \\citet{cohen:08} proposes that this is a manifestation of the mechanism described in \\citet{kilic:07} and \\citet{kalirai:07}, namely that metal-rich dwarfs lose enough mass to skip the horizontal branch resulting in a comparatively small fraction of metal-rich giants. Otherwise, if bulge dwarfs and giants share the same underlying MDF, either a systematic offset between dwarfs and giants or small number statistics could produce an artificially low observed probability. Preliminary metallicities of six bulge dwarfs from the 2009 bulge season foreshadow a potential shift toward lower metallicity in the bulge dwarf MDF, which would bring it into closer agreement with the bulge giants \\citep{bensbyIAU:09}.  Most elemental abundances confirm the trends observed in the other bulge dwarfs. The $\\alpha$-elements follow the bulge giants' trend and display solar abundance ratios befitting their high metallicity. This indicates that SNIa enriched the ISM by the time OGLE-2007-BLG-514S formed, making it one of the younger stars in the bulge. The bulge dwarfs provide the first large sample of iron-peak abundances in the bulge. Although we lack high metallicity disk data, we note that the bulge [V/Fe] rises at high [Fe/H].  In [Na/Fe], we find a much tighter relationship among bulge dwarfs than bulge giants at [Fe/H]$\\gtrsim 0$, despite claims by \\citet{lecureur:07} that their bulge giants reflect the composition of the ISM at formation. While the bulge dwarf [Na/Fe] measurements could in principle be in error, we consider this unlikely given their small scatter and robustness relative to the solar abundance. One possibility is that giants experience enhanced extra mixing only in the metallicity region not yet sampled by the bulge dwarfs (eg. $0.0\\lesssim\\mathrm{[Fe/H]}\\lesssim0.3$). Chemical evolution in the bulge could also produce both dwarfs and giants with high Na abundance in that metallicity region. Otherwise, enhanced extra mixing or measurement errors produced the giants' high [Na/Fe] and large scatter compared to the bulge dwarfs.  OGLE-2007-BLG-514S exhibits a strikingly low [Ba/Fe]$=-0.63\\pm0.29$ dex in comparison to most other bulge dwarfs. Ba is typically produced in the $s$-process, although the $r$-process does generate small amounts. We propose that the $r$-process dominates in the bulge like it does in the halo. Alternatively, metal-rich AGB stars can overproduce first $s$-process peak elements relative to Ba. Although Ba represented the only measurable neutron-capture element in this star, future analyses of $r$-process and first $s$-process peak elements within the bulge can constrain which mechanism is responsible for the low Ba abundance in this star.  Gathering a larger sample of bulge dwarfs will improve the dwarf MDF and fill out abundance trends. The study of this important stellar population is made possible by microlensing. Bulge dwarfs, like OGLE-2007-BLG-514S, can serve as a valuable tool for probing the formation and chemical evolution of the bulge."
    },
    "0910/0910.4261_arXiv.txt": {
        "abstract": "{We explore radio and spectroscopic properties of a sample of 14 miniature radio galaxies, i.e. early-type core galaxies hosting radio-loud AGN of extremely low radio power, 10$^{27-29}$ erg s$^{-1}$ Hz$^{-1}$ at 1.4 GHz.  Miniature radio galaxies smoothly extend the relationships found for the more powerful FR~I radio galaxies between emission line, optical and radio nuclear luminosities to lower levels.  However, they have a deficit of a factor of $\\sim$100 in extended radio emission with respect to that of the classical example of 3CR/FR~I. This is not due to their low luminosity, since we found radio galaxies of higher radio core power, similar to those of 3CR/FR~I, showing the same behavior, i.e. lacking significant extended radio emission.  Such sources form the bulk of the population of radio-loud AGN in the Sloan Digital Sky Survey. At a given level of nuclear emission, one can find radio sources with an extremely wide range, a factor of $\\gtrsim$100, of radio power.  We argue that the prevalence of sources with luminous extended radio structures in flux limited samples is due to a selection bias, since the inclusion of such objects is highly favored. The most studied catalogues of radio galaxies are thus composed by the minority of radio-loud AGN that meet the physical conditions required to form extended radio sources, while the bulk of the population is virtually unexplored. ",
        "introduction": "The presence of a radio source represents the most common manifestation of nuclear activity in early type galaxies. These sources are associated with a fraction as high as 40\\% of all bright elliptical and lenticular galaxies (e.g. \\citealt{sadler89,auriemma77}).  Due to the steepness of their luminosity function, most radio sources are of low luminosity. The interest in the properties of these objects lies in the fact that they sample an essentially unexplored regime for radio loud AGN (hereafter RLAGN) and they effectively close the gap between active and quiescent galaxies.  \\citet{balmaverde06b} considered a sample of low luminosity radio sources ($\\sim 10^{26-29}$ erg s$^{-1}$ Hz$^{-1}$ at 5 GHz) well suited for such a study. They are hosted by early-type galaxies selected on the basis of the presence of a shallow core in their host surface brightness profiles, defined as ``core'' galaxies (hereafter CoreG).  They used HST and Chandra data to isolate their optical and X-ray nuclear emission, showing that CoreG invariably host radio-loud nuclei, with an average radio loudness parameter of $R = L_{5\\rm {GHz}} / L_{\\rm B}$ $\\sim$ 4000, similar to the value measured for FR~I radio galaxies. Their optical and X-ray nuclear luminosities correlate with the radio-core power, smoothly extending the analogous correlations found for FR~I low luminosity radio galaxies \\citep{chiaberge:ccc,balmaverde06a} toward even lower power, by a factor of $\\sim$1000, covering a combined range of 6 orders of magnitude.  The similarities between CoreG and FR~I include the distributions of black hole masses, host galaxy luminosities, and also the properties of the surface brightness profiles \\citep{deruiter05}. This indicates that they are drawn from the same population of early-type galaxies.  The level of activity in these sources is closely related to the accretion rate of hot gas derived analyzing Chandra images.  \\citet{balmaverde08} considered a sample of 44 galaxies (CoreG and FR~I) and found that the accretion power correlates linearly with the jet power. These results strengthen and extend the validity of the results obtained by \\citet{allen06} and \\citet{hardcastle07}, indicating that hot gas accretion is the dominant process in powering FR~I radio galaxies across their full range of radio-luminosity. CoreG follow the same relationship, scaled to correspondingly lower accretion power.  This result, combined with the analogy of the nuclear properties, lead \\citet{balmaverde06b} to the conclusion that CoreG can be effectively considered as miniature radio galaxies.  The exploration of the properties of these sources offers the opportunity to probe the physical properties of RLAGN at widely different levels of activity with respect to classical FR~I radio galaxies. The aim of this paper is to expand the study of \\citet{balmaverde06b} by including i) an optical spectroscopic characterization of CoreG (derived by measuring ratios of emission lines) and ii) a detailed analysis of their radio properties (considering both morphologies and spectra), in order to perform a comparison with radio galaxies of higher power.  The main finding of this study is that while from the nuclear point of view CoreG are simply scaled down version of FR~I they have a deficit of a factor of $\\sim$100 in extended radio emission with respect to that of classical FR~I. We will show that, rather surprisingly, this is a general property, not limited to these faint radio sources, and that the vast majority of the population of RLAGN does not show the characteristic very luminous extended radio structures.  The paper is organized as follows. In Sect.  \\ref{sample} we present the sample selection. In Sect. \\ref{radiop} and \\ref{spectrop} we analyze the radio and optical spectroscopic properties of the CoreG. This leads to a multiwavelength view of their nuclei (Sect. \\ref{multip}) that is compared with those of the more powerful 3CR/FR~I. What emerges is that CoreG have a substantial deficit of extended radio emission. In Sect.  \\ref{bulk} we show, considering different samples of radio sources, that radio galaxies with feeble extended structures represent the bulk of radio-loud AGN population. In Sect.  \\ref{discussion} we discuss the consequences of this result on the general properties of radio galaxies and on our understanding of their nuclear activity. In Sect.  \\ref{summary} we summarize our findings and present our conclusions. In two Appendices we describe the details of the optical spectroscopic observations and explore the multiphase interstellar medium of CoreG. ",
        "conclusions": "\\label{summary}  We considered a sample of 14 luminous and nearby early-type galaxies hosting miniature RLAGN of extremely low radio luminosity (10$^{27-29}$ erg s$^{-1}$ Hz$^{-1}$ at 1.4 GHz) for which an extensive radio and optical analysis is performed. We collect the multi-wavelength nuclear information for our sample of core galaxies in order to compare them with those of the more powerful FR~I radio galaxies.  A radio analysis of these sources reveals that in many CoreG, the extended radio morphology is indicative of a collimated outflow. However, they are substantially less extended (up to $\\sim$ 20 kpc) and their core dominance is a factor of 40-100 higher than in more powerful FR~I radio sources in the 3CR sample. Their radio spectral properties are instead similar to those of 3CR/FR~I.  We also obtained new optical spectroscopic observations. CoreG show emission line ratios typical of AGN (they can be classified as low excitation galaxies), similar to, but of even lower excitation than, those of 3CR/FR~I radio galaxies, and matching more closely those of radio-quiet AGN of the same of line luminosity. This indicates that the line ratios are driven mostly by the AGN luminosity rather than by its radio-loudness.  The main result obtained from the spectroscopic study is that the observed line emission has generally an AGN origin and this enables us to include these emission lines in our analysis of the AGN properties. While miniature radio galaxies follow similar relationships to those found for more powerful radio galaxies in terms of line, optical, and radio nuclear luminosities, as well as accretion rate, they have a deficit of a factor of $\\sim$100 in extended radio emission with respect to that of classical 3CR/FR~I.  From the radio data alone, the possibility that this is due to a Doppler boosting enhancement of the radio core emission was still viable. The inclusion of emission lines (independent of orientation) leads us to conclude that what we are seeing in CoreG is a genuine deficit of extended radio emission, considering different estimators of nuclear activity, with respect to 3CR/FR~I.  The deficit of extended radio emission is also found in radio sources extracted from the B2 sample, with radio cores as powerful as those of classical FR~I radio galaxies. Therefore, the difficulty to produce prominent extended radio structures is not simply due to a low AGN luminosity. Actually, core dominated sources form the bulk of radio-loud AGN population in the SDSS/NVSS sample. The vast majority of these sources show, at a given level of line luminosity, a deficit of radio emission of a factor of $\\sim$ 100. Furthermore, this value should be considered as an upper limit due to the significant contribution of the radio core and since the samples considered are not completely free from selection biases.  At a given level of nuclear emission, one can find radio sources with an extremely wide range of radio power. Nonetheless, even the objects with the lowest level of extended radio emission are not radio-quiet. Not only are their nuclei radio-loud, but they have on average a ratio of radio to line luminosity larger by a factor of 300 with respect to radio-quiet AGN of similar radio power.  One possibility to explain this effect is to assume that this is driven by the source age, assuming that the radio luminosity increases with time as the radio source expands. A high number of young/faint sources is expected if the vast majority of RLAGN is short lived and never grows to large scales.  If this is the correct interpretation, the questions are: what mechanism sets the lifetime of the radio galaxies? Why are most radio-loud galaxies short-lived, while only a minority is active over long timescales?  The prevalence of low core dominance sources in flux-limited samples is apparently due to a selection bias, since the inclusion of sources with luminous extended radio structures is highly favored.  This result has several important ramifications. For example, unified models for RLAGN should be at least in part reconsidered, since the bulk of the RLAGN population is not well represented in the most studied samples of radio galaxies. Similarly, the relationship between line and radio emission results at least in part from selection effects, while apparently the link between the line and radio core is more robust and relies on the common dependence of these quantities on the strength of the nuclear continuum emission.  A thorough study of the morphology and colors of the hosts of core dominated sources, in particular those that are part of the large SDSS/NVSS sample, as well as of their radio morphology and optical spectroscopic classification is clearly required. This will enable us to perform a more detailed comparison with classical extended radio galaxies to reveal further similarities or differences that might account for their diversity.  The samples considered here are limited to AGN with $L_{\\rm [O III]} \\lesssim 10^{41}$ erg s$^{-1}$. It would be of great interest to explore whether the dominance of radio galaxies with relatively low extended radio emission applies also to sources of higher power.  \\appendix \\label{append}"
    },
    "0910/0910.4895.txt": {
        "abstract": "{}{}{}{}{} % 5 {} token are mandatory  \\abstract % context heading (optional) % {} leave it empty if necessary {SuperAGILE is the hard X-ray monitor of the AGILE gamma ray mission, in orbit since 23$^{rd}$ April 2007. It is an imaging experiment based on a set of four independent silicon strip detectors, equipped with one-dimensional coded masks, operating in the nominal energy range 18-60 keV.} % aims heading (mandatory) {The main goal of SuperAGILE is the  observation of cosmic sources simultaneously with the main gamma-ray AGILE experiment, the Gamma Ray Imaging Detector (GRID). Given its $\\sim$steradian-wide field of view and its $\\sim$15 mCrab day-sensitivity, SuperAGILE is also well suited for the long-term monitoring of Galactic compact objects and the detection of bright transients.} %methods {The SuperAGILE detector properties and design allow for a 6 arcmin angular resolution in each of the two independent orthogonal projections of the celestial coordinates. Photon by photon data are continuously available by the experiment telemetry, and are used to derive images and fluxes of individual sources, with integration times depending on the source intensity and position in the field of view.} % results heading (mandatory) {In this paper we report on the main scientific results achieved by SuperAGILE over its first two years in orbit, until April 2009. The scientific observations started on mid-July 2007, with the Science Verification Phase, continuing during the complete AGILE Cycle 1 and the first $\\sim$half of Cycle 2. Despite of the largely non-uniform sky coverage, due to the pointing strategy of the AGILE mission, a few tens of Galactic sources were monitored, sometimes for unprecedently long continuous periods, detecting also several bursts and outbursts. Approximately one gamma ray burst per month was detected and localized, allowing for prompt multi-wavelength observations. Few extragalactic sources in bright states were occasionally detected as well. The light curves of sources measured by SuperAGILE are made publicly available on the web in nearly realtime. For a proper scientific use of these we provide the reader with the relevant scientific and technical background. } % conclusions heading (optional), leave it empty if necessary {} ",
        "introduction": "The SuperAGILE experiment (Feroci et al. 2007) is the hard X-ray imager of the Italian gamma-ray astronomy mission AGILE (Tavani et al. 2009). The primary instrument of the AGILE payload is the GRID (Gamma Ray Imaging Detector, Barbiellini et al. 2002, composed of a Silicon Tracker (ST, Prest et al. 2003) and a small CsI(Tl) \"mini\"-calorimeter (MCAL, Labanti et al. 2009), surrounded by a plastic anticoincidence system (ACS, Perotti et al. 2006).  The main scientific goal of AGILE is the observation of the gamma-ray sky in the energy range from 50 MeV to a few GeV, over a field of view in excess of 2 steradian, ten years after the demise of the EGRET experiment onboard the Compton Gamma Ray Observatory. Several classes of celestial sources are known to emit gamma-rays in this energy band, most notably rotation powered pulsars, active galactic nuclei (especially blazars), supernova remnants and, occasionally, gamma-ray bursts. Other types of sources, such as galactic microquasars, low and high mass X-ray binaries or magnetars were expected to possibly emit gamma-rays in some special conditions (e.g., from shocks in the jet or in the accretion flow), but none of them had been unambiguously classified as a gamma ray source prior to the recent observations by AGILE or Fermi-GLAST (with the exception of the debated cases of LS 5039 and LSI +61 303).  The SuperAGILE experiment was designed to guarantee the hard X-ray (18-60 keV) monitoring in the tens of mCrab sensitivity range of the central $\\sim$steradian of the GRID field of view, aiming at identifying in the hard X-rays sources and/or conditions of special interest to the GRID, providing accurate ($\\sim$arcmin) localizations and allowing for multi-band observations. In addition to that, SuperAGILE independently acts as a wide field monitor for bright galactic sources, with photon by photon transmission and 5-$\\mu$s time accuracy.  A detailed description of the SuperAGILE experiment (pre-flight) may be found in Feroci et al. (2007), while information about the early in-flight operation and technical performance may be found in Feroci et al. (2008). Here we summarize the main instrument properties relevant to the aim of the present paper, that is providing an overview of the main scientific results achieved by the experiment during the first 20 months of observations from space, as well as guidelines and limits for a proper scientific use of the SuperAGILE data.  SuperAGILE is a coded-mask experiment devoted to image the transient hard X-ray sky by means of a set of four one-dimensional Silicon microstrip detectors, encoding the sky in two orthogonal directions. The field of view (FOV) of each pair of detectors is rectangular, 107$^{\\circ}$ x 68$^{\\circ}$ at zero response (Fig. \\ref{fig_fov}). Thus, the central 68$^{\\circ}$ x 68$^{\\circ}$ region is encoded in both directions (2 x 1D), while the peripheral $\\sim\\pm$20$^{\\circ}$ of the FOV are coded in one dimension only. Due to severe limitations at design and operation level (see Feroci et al. 2007, 2008a), SuperAGILE operates in the non-optimal (for a Silicon detector) energy range from 18 to 60 keV, being noise-limited at the low energy bound and efficiency limited at the upper bound, and behind a 5 mm thick shield of plastic scintillator (part of the AGILE ACS). The energy resolution is $\\sim$8 keV Full Width at Half Maximum, energy-independent. The one-dimensional point spread function (PSF) is 6 arcmin on-axis and the point source location accuracy reaches 1-2 arcmin for detections above $\\sim$10$\\sigma$, including a $\\sim$1 arcmin systematics contribution from the AGILE satellite attitude reconstruction. The SuperAGILE on axis 5-$\\sigma$ sensitivity is $\\sim$20 mCrab for a 50 ks net exposure (independently in each direction, that is half of the area of the experiment). Due to the triangular response of the collimator, the FOV at half sensitivity with 2 x 1D imaging is approximately 30$^{\\circ}$ x 30$^{\\circ}$. Instead, in each direction the sensitivity is better than 100 mCrab (5$\\sigma$, 50 ks, two detectors only) within the central 40$^{\\circ}$ x 80$^{\\circ}$ region (see Fig.\\ref{fig_fov}). The on-axis effective area is 250 cm$^{2}$ (total of 4 detectors, at 20 keV). The experiment background is dominated by the diffuse X-rays. Particle induced background is largely discriminated in amplitude and by the anticoincidence logic with the ACS, that actually causes a large dead-time effect when crossing the south atlantic anomaly, preventing in practice scientific observations in this orbital phase (see, e.g., Feroci et al. 2008a).   \\begin{figure} \\centering \\includegraphics[width=9cm, clip]{fig1.ps} \\caption{Representation of the field of view of the two pairs of SuperAGILE detectors, encoding two directions in the sky. Here we show with solid lines the sensitivity contours of one pair of detectors (i.e., one coding direction, either X or Z). The contours in dotted lines represent the sensitivity of the other pairs of detectors, rotated by 90$^{\\circ}$, providing the other coordinate. Sensitivity is given in mCrab (20-60 keV) at 5-$\\sigma$ confidence level, over a 50 ks integration time, for one direction only (half of the experiment area). } \\label{fig_fov} \\end{figure}   The SuperAGILE experiment entered its nominal observing phase on mid-July 2007, after the successful completion of its Commissioning Phase. Due to the inaccessibility of the Crab Nebula in June, caused by the  Sun constraints of the spacecraft, the first light of the experiment was taken on the Vela region, with the detection of the galactic binary system GX 301-2. Between July and October 2007 the experiment was calibrated in flight, using the Crab Nebula to scan several positions in the FOV, as a part of its Science Verification Phase. The latter was completed by the end of November 2007 with the observation of the Vela field, the Galactic Center and the Cygnus region. On December 2007 the AGILE Observing Cycle 1 started, with typical continuous exposures of 4 weeks to the same field. Contrary to those of the GRID, the SuperAGILE data of Cycle 1 are not open to Guest Observers. However, light curves of the sources detected by SuperAGILE are in public distribution in nearly realtime through a web page (see $\\S \\ref{sa_web}$).  In the following sections we will describe the scientific data products of SuperAGILE ($\\S$ \\ref{software}), review the status of experiment calibration and pointing strategy ($\\S$ \\ref{calib} and $\\S$ \\ref{pointings}), and overview the main scientific observations carried out over the first $\\sim$20 months of nominal operations in orbit ($\\S$ \\ref{science}), including a preliminary list of the detected sources. ",
        "conclusions": "The SuperAGILE experiment is successfully operating onboard the AGILE mission since 2007, April 23$^{rd}$. In this paper we provided a description of the SA data, and reported an overview of the main scientific results achieved in the first $\\sim$20 months of scientific observations. The main goals of SuperAGILE are the simultaneous hard X-ray observation of the central region of the field of view of the AGILE/GRID experiment, the prompt discovery and localization of gamma-ray bursts and Galactic transients, and the long-term monitoring of the bright Galactic sources.  The aim of the simultaneous hard X and gamma-ray observation of the same field is to discover correlated variability of sources in these two energy ranges, and use the finer angular resolution of SuperAGILE to localize gamma-ray sources detected in the GRID. This is the first time that an X-ray and a gamma-ray imager systematically observe the same field simultaneously. Regrettably, this did not bring yet to discover any previously unknown correlated behavior of sources. The sources that SuperAGILE and GRID detected simultaneously were already known to emit in both energy ranges. This is the case for the Crab pulsar, or the AGNs 3C 273 and Mkn 421. However, the guaranteed simultaneous GRID and SuperAGILE observations represent the basic seed of multifrequency campaigns, often complemented with radio to TeV observations. This led, for example, to detect a flaring state of the BL Lac source Mkn 421 and interpret the time variability of the simultaneous optical-to-TeV spectral energy distribution in terms of a rapid acceleration episode of the leptons in the jet (Donnarumma et al. 2009).  The primary goal is of course to search for positive detections in both instruments, but there are cases where also upper limits in one or the other are important to interpret the origin of the detected emission. This was the case of the unidentified transient sources discovered by the AGILE/GRID on the galactic plane, in the Cygnus and Musca regions (e.g., Longo et al. 2008; Pittori et al. 2008; Chen et al. 2008; Giuliani et al. 2008b). Following Romero \\& Vila (2009), the SuperAGILE upper limits on the simultaneous hard X-ray emission of these transient gamma-ray sources imply a very large ratio between the gamma-ray and X-ray luminosity, posing severe constraints on their interpretation in terms of emission from jets of galactic microquasars, strongly favoring a dominant hadronic component in the jet.  Interesting results were obtained also in the field of gamma-ray bursts. Despite of the prompt SuperAGILE arcmin-localization of several GRBs, enabling searches in time and spatial coincidence, in the reported period only two events were significantly detected by the AGILE/GRID (Giuliani et al. 2008a, Moretti et al. 2009), including GRB080514B, the first GRB ever for which the gamma-ray emission can be correlated to an afterglow counterpart with a measured distance, and GRB 090401B. The detection of these events by all the AGILE instruments - SA (18-60 keV), GRID ($>$25 MeV), MCAL (350-700 keV) and ACS ($>$8 keV) - also allowed to measure a few seconds delay of the emission above 25 MeV with respect to that at lower energies. The observations of the \\textit{Compton Gamma Ray Observatory}, particularly the EGRET and BATSE experiments, provided the result that the GRBs with a gamma-ray counterpart detectable above the EGRET sensitivity (that is similar to AGILE/GRID's) belong to the brightest 5\\% of the BATSE GRB distribution, but the statistics was only that of 5 events. Although the SuperAGILE GRBs are selected towards the bright end of the BATSE distribution, because of its smaller area, the 2 GRID detections over 21 GRBs localized by SuperAGILE by the end of April 2009 are still consistent with the approximate EGRET statistics.   The pointing strategy of the AGILE mission, typically consisting of long exposures, offers as a by-product a long term monitoring of the same field by SuperAGILE and as a drawback a very inhomogeneous sky coverage (see Fig. \\ref{fig_exposure}), largely privileging the Cygnus and Vela sky regions in particular. Although the analysis of the SuperAGILE data collected so far is not complete yet (see $\\S$\\ref{software} and $\\S$\\ref{detected_sources}), about 60 sources were detected over the time period reported in this paper. As it can be seen from Table \\ref{table:list}, the large majority of them are X-ray binaries, mostly Low Mass X-ray Binaries (LMXB).  Despite the inhomogeneity and incompleteness of the SuperAGILE exposure and data analysis, it may be useful to put the list of sources detected by SA in the context of the current scenario in the hard X-rays. Currently, INTEGRAL/ISGRI and Swift/BAT offer the most complete sky surveys in the hard X-rays. The list of sources detected by SuperAGILE given in Table \\ref{table:list} was obtained mainly by daily integrations. The sensitivity of SuperAGILE on this timescale is in the range of $\\sim$15-20 mCrab. Applying a flux cut at 15 mCrab in the ISGRI 3.5-year catalogue (20-40 keV, Bird et al. 2007) and to the BAT 22-month catalog (14-150 keV, Tueller et al. 2009) we found about 50 and 40 sources, respectively. Of these, in ISGRI 38\\% are HMXB and 57\\% are LMXB, while in BAT they are 31\\% and 59\\%, respectively. From the list in Table \\ref{table:list}, SA detected 57 persistent sources, of which 33\\% are HMXB and 58\\% are LMXB. As expected, a large number of sources are actually in common, being the bright tail of the LogN-LogS distribution of persistent hard X-ray sources. Although the agreement is remarkably good, the above comparison should not be taken rigorously, due to a large number of important caveats about pointing strategy, type of the analysis, exposure, energy range and intrinsic source variability. However, it is a nice general confirmation that the source detection record by SuperAGILE is consistent with the sensitivity limit of the current analysis. \\\\ Pushing the comparison (as well as the caveats) beyond, we made the same type of selection to the source catalog of the BeppoSAX/Wide Field Cameras (WFC, Verrecchia et al. 2007). This experiment operated in the 2-26 keV energy range. Selecting the sources with reported average flux above 15 mCrab in 2-10 keV, we found $\\sim$80 sources, of which 20\\% are HMXB and 73\\% are LMXB. A comparison between the type and number of X-ray binary sources brighter than 15 mCrab below $\\sim$20 keV (in the WFC) and above (in SA, BAT and ISGRI) shows that approximately the same number of objects is found among the HMXB, while the LMXB decrease to one half above 20 keV. This suggests that the LMXB, as a class, tend to have energy spectra softer than the Crab (photon index $\\sim$2.1), while the HMXB tend to have harder spectra, on average.   Many of the SA sources were detected on the orbital timescale, meaning that public light curves are available for them at the ASDC web page (see $\\S$\\ref{sa_web}), updated daily with new data (when available). In this paper we provided detailed explanations on how these light curves are extracted and under what assumptions, thus offering an understanding of the limits of their scientific use. Other sources have been detected either in longer integrations or in short outbursts. For these sources no publicly accessible science products are currently available yet.  The main characteristics of the SA observation is the continuity and long duration. Some galactic sources have been observed for probably the longest continuous stretches of time ever, allowing for monitoring the long term variability of their hard X-ray flux. The duration and continuity of the monitoring distinguishes the SA data from those of more sensitive experiments (e.g., Swift/BAT) that offer short and sparse observations of individual sources, sometimes missing short-term variability or events. In addition, the SA field is always simultaneously observed by the AGILE/GRID, thus offering the simultaneous measurement of the flux above $\\sim$50-100 MeV.       The SuperAGILE experiment is operating nominally since its switch-on in orbit in 2007, showing no signs of degradation or losses with time. %, just like the rest of the payload and the %engineering systems. The in-flight operation of the AGILE mission is expected to continue until at least mid-2011, subject to approval by the Italian Space Agency (ASI)."
    },
    "0910/0910.4866_arXiv.txt": {
        "abstract": "The antenna array LOPES is set up at the location of the KASCADE-Grande extensive air shower experiment in Karlsruhe, Germany and aims to measure and investigate radio pulses from Extensive Air Showers. The coincident measurements allow us to reconstruct the electric field strength at observation level in dependence of general EAS parameters. In the present work, the lateral distribution of the radio signal in air showers is studied in detail. It is found that the lateral distributions of the electric field strengths in individual EAS can be described by an exponential function. For about 20\\% of the events a flattening towards the shower axis is observed, preferentially for showers with large inclination angle. The estimated scale parameters $R_0$, describing the slope of the lateral profiles range between $100$ and $200\\,$m. No evidence for a direct correlation of $R_0$ with shower parameters like azimuth angle, geomagnetic angle, or primary energy can be found. This indicates that the lateral profile is an intrinsic property of the radio emission during the shower development which makes the radio detection technique suitable for large scale applications. ",
        "introduction": "The traditional method to study extensive air showers (EAS), which are generated by high-energy cosmic rays entering the Earth's atmosphere, is to measure the secondary particles with sufficiently large particle detector arrays. In general, these measurements provide only immediate information on the status of the air shower cascade at the particular observation level. This hampers the determination of the properties of the primary inducing the EAS as compared to methods like the observation of Cherenkov and fluorescence light, which also provide information on the longitudinal EAS development, thus providing a more reliable access to the information of interest~\\cite{rpp}.  In order to reduce the statistical and systematic uncertainties of the detection and reconstruction of EAS, especially with respect to the detection of cosmic particles of highest energies, measuring the radio emission during the shower development is being discusses as a new detection technique. The radio waves can be recorded day and night, and provide a bolometric measure of the electromagnetic shower component. Due to technical restrictions in past times, the radio emission accompanying cosmic ray air showers was a somewhat neglected feature in the past. However, the study of this EAS component has experienced a revival by recent activities, in particular by the LOPES project~\\cite{Falcke05,Horn06}. LOPES as pathfinder for large scale radio detection for the LOFAR project~\\cite{lofar} and the Pierre Auger Observatory~\\cite{auger} investigates the correlations of radio data with shower parameters reconstructed by the extensive air shower experiment KASCADE-Grande~\\cite{navarra}. Hence, LOPES, which is designed as a digital radio interferometer using large bandwidth and fast data processing, profits from the reconstructed air shower observables of KASCADE-Grande.  The main goal of the investigations in the frame of LOPES is the understanding of the shower radio emission in the primary energy range of $10^{16}\\,$eV to $10^{18}\\,$eV. Of particular interest is the investigation of the correlation of the measured field strength with main shower parameters. These are the orientation of the shower axis (azimuth angle, zenith angle, and the geomagnetic angle, i.e.~the angle between shower axis and geomagnetic field), the position of the observer relative to the shower axis, and the energy and mass (electron and muon number) of the primary particle. Another goal of LOPES (LOPES$^{\\rm STAR}$) is the optimization of the hardware (antenna design and electronics) for a large scale application of the detection technique including a self-trigger mechanism for a stand-alone radio operation~\\cite{Asch07}.  In the present study we investigate in detail the lateral profile of the radio signal as measured by LOPES. Due to a precise amplitude calibration of each individual antenna and the event information from KASCADE-Grande, this is possible on an event-by-event basis with high accuracy. Such investigations are of great interest as the lateral shape defines the optimum grid size for a radio antenna array in a stand-alone mode. Of particular interest is the scale parameter which describes the amount of the signal decrease with distance from the shower axis and the dependence of that parameter on characteristics of the primary particle. In addition, knowing the lateral extension in detail will contribute to the understanding of the emission mechanism of the radio signal as simulations have shown~\\cite{Huege08}, that the lateral shape can be related to important physical quantities such as the primary energy or the mass of the primary. ",
        "conclusions": "The described studies on lateral distributions of the radio signal in air showers are based on coincidence measurements of the LOPES radio antenna array and the air shower experiment KASCADE-Grande. This is a unique combination of two different detection techniques for EAS investigations. Furthermore, the obtained data of both experiments are well calibrated and can therefore be used to compare the properties of the radio emission with EAS parameters in detail. Such investigations aim to fully understand the radio emission in extensive air showers for primary energies below $10^{18}\\,$eV. In the frame of this work a calibrated data set of showers has been analyzed to study the lateral distribution of the radio emission in individual air showers.  The applied analysis method uses up-sampled signals from individual antennas to derive the electric field strength per unit bandwidth $\\epsilon$ (40-80$\\,$MHz and east-west polarization cmponent, only). With the help of the reconstructed shower parameters from KASCADE-Grande the lateral distribution of the radio signal in EAS for individual events is obtained. Here, the uncertainty of the amplitude calibration and an antenna-by-antenna event-by-event background calculation enters into the uncertainty of the determined field strength $\\epsilon$. The uncertainty of the reconstructed geometry of the shower contributes to the uncertainty for the lateral distance $R$ of each antenna. For 110 showers which have a large enough radio signal in all antennas the lateral shape of the signal as well as correlations with global shower parameters, like direction and energy, were studied.  Exponential as well as power law functions were applied to the measured profiles, where the first one slightly better describes the data. The exponential function has two free parameters, the field strength $\\epsilon_0$ at the shower axis, and the scale parameter $R_0$. The scale parameter $R_0$ is a quantity that describes the slope of the lateral profile. The scale parameter distribution shows a peak value of $R_0\\approx125$~m and has a tail with very flat lateral distributions. Excluding the flat events a mean value of $\\bar{R_0} \\approx 150\\,$m with a width of $\\sigma=50\\,$m was obtained. No direct evidence for a dependence on the shower parameters azimuth angle, geomagnetic angle, and primary energy could be found. This indicates that the lateral profile is an intrinsic property of the radio emission and the shower development. Comparing the obtained scale parameter with published values of earlier experiments, a good agreement has been found. Studying the lateral distributions in individual events, approximately 20\\% of the studied showers show a very flat lateral distribution or exhibit a flattening towards the shower center. Preferably, such showers arrive under larger zenith angle and axes are close to the antennas. But, there are too few such showers measured to derive any significant correlation with specific shower parameters. However, the result is an indication that there might be (at least for a part of the showers) additional, not yet understood effects of the radio emission during the shower development. In the near future, besides higher statistics delivered by LOPES, the LOFAR project with its hundreds of antennas will provide detailed information (including polarization effects) on the lateral distribution of the radio signal in EAS.  The present analysis concentrates on experimental findings, only, but for the understanding of the radio emission observed in EAS a comparison with detailed Monte-Carlo simulations on an event-to-event basis will be very important. This will be possible soon because of the performed amplitude calibration of LOPES and the estimate of the field strength at individual antennas. Due to a new simulation strategy, using more realistic air shower models derived from per-shower CORSIKA~\\cite{cors} simulations including radio emission simulations with REAS2~\\cite{Huege07} and a simulation of the antenna response, detailed comparisons between LOPES measured events and simulations can be performed.  The future goal of the radio detection technique lies in a large scale application with hundreds of antennas. The quadratic dependence of the received radio power on primary energy and the comparatively (with respect to the electromagnetic component of EAS) large lateral scale parameter will make radio detection a cost effective method for measuring the longitudinal development of air showers of the highest energy cosmic rays.  \\begin{ack} LOPES has been supported by the German Federal Ministry of Education and Research. The KASCADE-Grande experiment is supported by the German Federal Ministry of Education and Research, the MIUR and INAF of Italy, the Polish Ministry of Science and Higher Education and the Romanian Ministry of Education and Research. \\end{ack}"
    },
    "0910/0910.3731_arXiv.txt": {
        "abstract": "We conducted a comprehensive study on the shell structure of the \\object{Cygnus Loop} using 41 observation data obtained by the \\textit{Suzaku} and the \\textit{XMM-Newton} satellites. To investigate the detailed plasma structure of the Cygnus Loop, we divided our fields of view into 1042 box regions. From the spectral analysis, the spectra obtained from the limb of the Loop are well fitted by the single-component non-equilibrium ionization plasma model. On the other hand, the spectra obtained from the inner regions are well fitted by the two-component model. As a result, we confirmed that the low-temperature and the high-temperature components originated from the surrounding interstellar matter (ISM)  and the ejecta of the Loop, respectively.  From the best-fit results, we showed a flux distribution of the ISM component. The distribution clearly shows the limb-brightening structure, and we found out some low-flux regions. Among them, the south blowout region has the lowest flux. We also found other large low-flux regions at slightly west and the northeast  from the center. We estimated the former thin shell region to be $\\sim1\\arcdeg.3$ in diameter and concluded that there exists a blowout along the line of sight in addition to the south blowout. We also calculated the emission measure distribution of the ISM component and showed that the \\object{Cygnus Loop} is far from the result obtained by a simple Sedov evolution model. From the results, we support that the \\object{Cygnus Loop} originated from a cavity explosion. The emission measure distribution also suggests that the cavity-wall density is higher in the northeast than that in the southwest. These results suggest that the thickness of the cavity wall surrounding the \\object{Cygnus Loop} is not uniform. ",
        "introduction": "\\label{sec:intro} The supernova (SN) explosions blow out the various heavy elements generated by the nucleosynthesis process inside the progenitor stars. Meanwhile, the blast wave generated by the SN explosion sweeps up and heats the interstellar matter (ISM). It forms a characteristic structure of each supernova remnant (SNR). In this way, the morphology of the shell structure provides information about the ambient density of the ISM.  The \\object{Cygnus Loop} is one of the brightest SNRs in the X-ray sky. Its age is estimated to be $\\sim10,000$yr and the distance is comparatively close to us (540pc; Blair et al. 2005). The large apparent size (2.5\\arcdeg$\\times$3.5\\arcdeg; Levenson et al. 1997) enables us to study the plasma structure of the Loop. The origin of the \\object{Cygnus Loop} is thought to be a cavity explosion \\citep{McCray79, Hester86, Hester94, Levenson97}. The progenitor is presumed to be a B0, 15\\MO \\ star \\citep{Levenson98} and some X-ray studies also estimated the progenitor mass to be 12-15\\MO \\ (\\textit{e.g.}, Tsunemi et al. 2007). From the morphological point of view, the \\object{Cygnus Loop} is a typical shell-like SNR and it is almost circular in shape. However, a large break is seen in the south, known as ``blowout'' region \\citep{Aschenbach99}. The origin of the ``blowout'' is not well understood. \\citet{Aschenbach99} explained this extended structure as a breakout into a lower density ISM. On the other hand, based on a radio observation, \\citet{Uyaniker02} suggested the existence of a secondary SNR in the south. Some other radio observations support this conclusion \\citep{Uyaniker04, Sun06}. Recently, \\citet{Uchida08} observed this region with the \\textit{XMM-Newton} observatory and found that the X-ray spectra of this region consist of two plasma components with different electron temperatures. Judging from the plasma structures and the metal distributions, they concluded that the X-ray emission is consistent with a \\object{Cygnus Loop} origin and that each plasma component is derived from the ejecta and the cavity material, respectively. They also showed that the X-ray shell is thin in their fields of view (FOV) and concluded that the origin of the blowout can be explained as a breakout into a lower density ISM as proposed by \\citet{Aschenbach99}.  It is natural to consider that such a breakout may also exist along the line of sight. \\citet{Tsunemi07} observed the \\object{Cygnus Loop} along the diameter with \\textit{XMM-Newton} and argued about it. They divided their FOV into a north path and a south path and found the thin shell region to be 5\\arcmin \\ in the south path and 20\\arcmin \\ in the north path. They estimated this thin shell region to have a diameter of 1\\arcdeg, centering on $\\alpha = 20^{\\mathrm h}49^{\\mathrm m}11^{\\mathrm s}$, $\\delta = 31\\arcdeg05\\arcmin20\\arcsec$. \\citet{Kimura09} expanded their observation northward with \\textit{Suzaku} and found that the flux of the swept-up matter in southwest is about a third of that in northeast. The width of this region is $\\sim50\\arcmin$. \\citet{Kimura09} presumed that there is a blowout region in the direction of our line of sight in the middle west of the Loop.  In this paper, we used 41 observation data obtained by the \\textit{Suzaku} and the \\textit{XMM-Newton} observatories. We reanalyzed all the data to conduct a comprehensive study on the shell structure of the \\object{Cygnus Loop}. ",
        "conclusions": "By analyzing the X-ray spectra, we clearly distinguished the ISM component from the ejecta component, and established a method to investigate the line-of-sight shell structure. From the flux distribution of the ISM component, we found three low-flux regions in the FOV; one is a well-known southwest blowout which is evidence of the cavity-wall break, and we also found other low-flux regions  at the west and the northeast of the \\object{Cygnus Loop} center. From the EM distribution of the ISM component, we support that the \\object{Cygnus Loop} is originated from a cavity explosion. Then, the ISM component, or cavity wall does not have an uniform structure but has a lot of breaks or tenuous regions. We also found that the condition of the surrounding cavity wall is not uniform;  the density of it is globally higher in the northeast than that in the southwest."
    },
    "0910/0910.4327_arXiv.txt": {
        "abstract": "We present an analysis of neutrinos detected with the Baikal neutrino telescope NT200 for correlations with gamma-ray bursts (GRB). No neutrino events correlated with GRB were observed. Assuming a Waxman-Bahcall spectrum, a neutrino flux upper limit of {\\bf $E^2 \\Phi < 1.1 \\times 10^{-6}cm^{-2}s^{-1}sr^{-1}GeV$ } was obtained. We also present the Green's Function fluence limit for this search, which extends two orders of magnitude beyond the energy range of the Super-Kamiokande limit. ",
        "introduction": "The Baikal neutrino telescope NT200 \\cite{1,2} is operating in Lake Baikal, Siberia, at a depth 1.1 km since April, 1998. NT200 consists of 8 strings of 70 m length: 7 peripheral strings and a central one. Interstring distances are about 20 m. Each string includes 24 pairwise arranged optical modules (OM). Each OM contains a 37-cm diameter hybrid photodetector QUASAR-370.  A number of relevant physics results has been obtained so far with the NT200 telescope, e.g. limits on the diffuse flux of extraterrestrial high energy neutrinos, limits on neutrino fluxes from Dark Matter annihilation (Sun, Earth), and on the flux of relativistic and slow magnetic monopoles \\cite{2,3,4}.  This work is devoted to the search of neutrino events correlated with observations of more than 300 gamma-ray bursts (GRBs) reported from 1998 to 2000 by the Burst and Transient Source Experiment (BATSE) \\cite{5}.  The detection strategy for neutrino events with the NT200 telescope is based on a search for Cherenkov light from relativistic up-going muons produced by neutrino interactions.  Information about the GRB time and location on the sky allows to reduce the atmospheric muon background and, as a result, significantly increases the sensitivity of the neutrino telescope to neutrino events correlated with GRB. ",
        "conclusions": "We have presented results of a search for neutrino induced muons detected with Baikal Telescope NT200 in coincidence with 303 gamma-ray bursts recorded from 1998 to 2000 by BATSE. NT200's field of view covers most part of the Southern hemisphere. No evidence for neutrino-induced  muons from gamma-ray bursts is found. The resulting Green's Function fluence limit for this search extends that of Super-Kamiokande by two orders of magnitude in energy. Assuming a Waxman-Bahcall spectrum, a neutrino flux upper limit of \\\\$E_{\\nu}^2 \\, \\Phi_{\\nu} \\leq 1.1 \\times 10^{-6} GeVcm^{-2}s^{-1}sr^{-1}$ is obtained."
    },
    "0910/0910.5860_arXiv.txt": {
        "abstract": "We present the two- and three-point real space correlation functions of the five-year {\\it WMAP} sky maps and compare the observed functions to simulated $\\Lambda$CDM concordance model ensembles. In agreement with previously published results, we find that the temperature correlation functions are consistent with expectations. However, the pure polarization correlation functions are acceptable only for the 33GHz band map; the 41, 61, and 94 GHz band correlation functions all exhibit significant large-scale excess structures. Further, these excess structures very closely match the correlation functions of the two (synchrotron and dust) foreground templates used to correct the {\\it WMAP} data for galactic contamination, with a cross-correlation statistically significant at the $2\\sigma-3\\sigma$ confidence level. The correlation is slightly stronger with respect to the thermal dust template than with the synchrotron template. ",
        "introduction": "Observations of the cosmic microwave background (CMB) have been among the most important ingredients in the revolution of cosmology that has taken place in the last two decades, when cosmology changed from a data starved to a data rich science. The so far most influential observations have been made with the {\\it Wilkinson Microwave Anisotropy Probe} ({\\it WMAP}; Bennett et al. 2003a; Hinshaw et al. 2007, 2009) satellite experiment, which has measured the microwave sky at five different frequencies (23, 33, 41, 61, and 94 GHz) in both temperature and polarization. Its main results are five sky maps with resolutions between $13'$ and $55'$, each comprising more than three million pixels in each of the three Stokes parameters I, Q, and U.    The detector technology leading to instruments like {\\it WMAP} has thus been very important in observing the CMB. Almost equally important has been the progress in computer technology and algorithms. With the enormously increased number of pixels from new high-resolution instruments, scientists have had to develop clever algorithms for every step of the required data analysis pipeline: map making (e.g., Ashdown et al. 2007; Hinshaw et al. 2003b, and references therein), component separation (e.g., Bennett et al. 2003b; Leach et al. 2008; Eriksen et al. 2008, and references therein), power spectrum and cosmological parameter estimation (e.g., G{\\'o}rski 1994; Lewis \\& Bridle 2002; Hivon et al. 2002; Hinshaw et al. 2003a; Verde et al. 2003; Eriksen et al. 2004d, and references therein), and analysis of higher-order statistics (e.g., Hinshaw et al. 1994; Kogut et al. 1995; Komatsu et al. 2003; Eriksen et al. 2004c, 2005 and references therein).  In this paper, we revisit a well-known example of the latter category, namely, real-space $N$-point correlation functions. These functions arise naturally in studies of non-Gaussianity, since it can be shown that any odd-ordered $N$-point function has an identically vanishing expectation value, while all even-ordered $N$-point functions have an expectation value given by products of the corresponding two-point functions. Violation of either of these relations would indicate the presence of a non-Gaussian component in the field under consideration.  Earlier $N$-point correlation function analyses of CMB data include studies of the {\\it COBE}-DMR data \\citep{hinshaw:1995, kogut:1996, eriksen:2002}, the {\\it WMAP} temperature data \\citep{eriksen:2005} and one single two-point analysis of the first-year {\\it WMAP} polarization data \\citep{kogut:2003}. This paper is the first to consider the much more mature five-year {\\it WMAP} polarization data, and the first to compute the three-point correlation function from any CMB polarization data. To do this, we adopt and extend the $N$-point algorithms developed by \\citet{eriksen:2004b}. ",
        "conclusions": "\\label{sec:conclusions}  We compute for the first time the polarized two- and three-point correlation functions from the five-year {\\it WMAP} data, with the initial motivation of constraining possible non-Gaussian signals on large scales in the CMB polarization sector. However, our main result is a detection of a signature consistent with residual foregrounds in both the $Q$- and $V$-band data, significant at $2\\sigma$-$3\\sigma$.  This should not come as a complete surprise, though, as a similar result was found in the three-year {\\it WMAP} data through a direct template fit approach by \\citet{eriksen:2007}. They used a Gibbs sampling implementation to estimate the joint CMB power spectrum and foreground template posterior, and found non-zero template amplitudes at similar significance levels when marginalizing over the CMB spectrum. Given the results found in the present paper, this still appears to be an important issue for the five-year {\\it WMAP} data release. Resolving this question is of major importance for the CMB community, since many cosmological parameters depend critically on the large-scale {\\it WMAP} polarization data, most notably the optical depth of reionization, $\\tau$, and the spectral index of scalar perturbations, $n_{\\textrm{s}}$.  The question of primordial non-Gaussianity in CMB polarization data can not be meaningfully addressed as long as this issue is unresolved. The three-point correlation function results shown in this paper therefore mostly serve as a demonstration of the approach, and as a cross-check on the two-point results. The three-point function will obviously become a more important and independent quantity in the future, when higher-fidelity polarization data become available."
    },
    "0910/0910.5267_arXiv.txt": {
        "abstract": "We report observations of 15 high redshift ($z = 1-5$) galaxies at 350$\\mu$m using the Caltech Submillimeter Observatory and SHARC-II array detector. Emission was detected from eight galaxies, for which far-infrared luminosities, star formation rates, total dust masses, and minimum source size estimates are derived.  These galaxies have star formation rates and star formation efficiencies comparable to other high redshift molecular emission line galaxies.  The results are used to test the idea that star formation in these galaxies occurs in a large number of basic units, the units being similar to star-forming clumps in the Milky Way. The luminosity of these extreme galaxies can be reproduced in a simple model with (0.9-30)\\ee6 dense clumps, each with a luminosity of 5\\ee5 \\lsun, the mean value for such clumps in the Milky Way. Radiative transfer models of such clumps can provide reasonable matches to the overall SEDs of the galaxies. They indicate that the individual clumps are quite opaque in the far-infrared. Luminosity to mass ratios vary over two orders of magnitude, correlating strongly with the dust temperature derived from simple fits to the SED. The gas masses derived from the dust modeling are in remarkable agreement with those from CO luminosities, suggesting that the assumptions going into both calculations are reasonable. ",
        "introduction": "\\label{intro}  Studies of star formation in the early Universe are important to an understanding of galaxy formation and evolution.  Observations of quasar host galaxies and sub-millimeter galaxies with detectable molecular line emission offer an opportunity to extend such studies to high redshift ($z>2$), albeit with a limited sample \\citep[see][]{annrev05}.  Following that review, we refer to these galaxies collectively as EMGs (Early Universe Molecular Emission Line Galaxies). Several kinds of models of these objects have been proposed. Dust in quasar host galaxies could be heated by the radiation from the central AGN (e.g., \\citealt{granato96}, \\citealt{andreani99}). While it is hard to rule out such models, general arguments tend to favor starbursts as the main power source (e.g., \\citealt{blain02}). In addition, detailed studies of well-known sources have indicated that star formation dominates the contribution from the black hole (\\citealt{rowan-robinson2000}, \\citealt{weiss03}, \\citealt{chapman04}, \\citealt{weiss07}).  There are also variations among models based on star formation. \\citet{err03} proposed a model in which the far-infrared emission arises from cirrus (relatively diffuse dust) heated by ultraviolet photons leaking out from star formation regions. However, most recent analysts have focused on models in which the far-infrared radiation arises from dust that is intimately associated with a burst of star formation (e.g., \\citealt{narayanan09a}). In the embedded starburst models, the total luminosity of the dust continuum emission is a measure of the star formation rate, and the luminosity of the molecular line emission or the dust emission at long rest wavelengths measures the amount of material available for star formation.  Here we report observations of the dust continuum emission at 350$\\mu$m wavelength, obtained at the Caltech Submillimeter Observatory\\footnote{The Caltech Submillimeter Observatory is supported by the NSF.}. This observed wavelength falls roughly near the peak of the emission in the rest frame of the objects observed and is, therefore, a desirable wavelength for observations to determine far-infrared luminosities. It will also provide stronger constraints on the characteristic temperature of the far-infrared radiation, which can distinguish between cirrus models and models of embedded star formation, as the former predict cool ($\\td < 30$ K) dust \\citep{err03}.  This work extends previous studies at 350$\\mu$m of high redshift molecular line emission galaxies \\citep{benford99, weiss03, beelen06, wang08sharc2, wang08cont}.  Throughout this paper we have assumed a $\\Lambda$CDM cosmology with $H_0 = 70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_m = 0.27$, and $\\Omega_{\\Lambda} = 0.73$ \\citep{spergel07}. ",
        "conclusions": "\\label{sourcedisc}  \\subsection{Detections}   \\it LBQS0018 \\rm is an optically identified quasar from the survey of \\citet{foltz89} detected in CO(3--2) with the IRAM Interferometer \\citep{izaak04}.  LBQS0018 is radio quiet.  \\it SMM J02396 \\rm was detected in CO(3--2) emission by \\citet{greve05} using the IRAM Interferometer. The source is a ring galaxy identified by \\citet{soucail99} from HST images. They measured the redshift of $z = 1.062$ and determined the magnification factor by the cluster Abell 370 to be $\\mu = 2.5$. \\citet{smail02} refer to the source as J02399-0134, rather than J02396-0134 as it has come to be labeled.  \\it RX J0911.4 \\rm is a ROSAT source identified as a mini-broad absorption line quasar by \\citet{bade97}.  The quasar is strongly lensed ($\\mu=22$) by a galaxy at z = 0.8 \\citep{kneib00}. RX0911.4 was detected in CO(3--2) at the Owens Valley Radio Observatory by \\citet{hainline04}.  The quasar is radio quiet.  \\it SMM J14011 \\rm was the second SCUBA source to be detected in a molecular emission line, namely CO(3--2) at the Owens Valley Radio Observatory \\citep{frayer99}.  It has been intensively studied since then in CO lines \\citep[see] [] {downes03}; the CI(1--0) line has been detected \\citep{weiss05}. The source is lensed, but the lens magnification is very uncertain and model dependent, ranging from $\\mu = 5 - 25$ \\citep{downes03}.  \\subsection{Tentative Detections}   \\it SMM J02399 \\rm is a hyperluminous infrared galaxy \\citep{ivison98}, the first SCUBA source to be detected in molecular line emission, by \\citet{frayer99} at the Owens Valley Radio Observatory in CO(3--2).  From higher angular resolution observations with the IRAM Interferometer, \\citet{genzel03} argued for a model with a very large ($r \\geq 2$ kpc) disk, a possibility that awaits still higher angular resolution confirmation. Our detection at 350 \\micron\\ of $29\\pm9$ mJy is somewhat inconsistent with the high value at 450 \\micron\\ \\citep{ivison98}. (Fig. 3), even considering the large errors in each.  \\it SDSS 0338 \\rm is a high-redshift ($z = 5.03$) quasar discovered by \\citet{fan99}.  The CO(5--4) line was detected with the IRAM Interferometer by \\citet{maiolino07}.  The radio continuum has not been detected.  This source was previously detected at 350 \\micron\\ by \\citet{wang08sharc2}, with  $S_{350\\mu m} = 17.7 \\pm 4.4$ mJy. Our value of $29\\pm10$ mJy is consistent within the uncertainties.  \\it MG0751 \\rm is a quasar detected in CO(4--3) emission by \\citet{barvainis02a} using the IRAM Interferometer.  The quasar is highly lensed, with magnification of 16.6 \\citep{barvainis02b}. This is our weakest detection, with a S/N ratio of only 2.3, weaker than the 450 \\micron\\ detection.  \\it SMM J09431 \\rm is a submillimeter galaxy detected in CO(4--3) emission by \\citet{neri03} using the IRAM Interferometer.  \\label{discussion}  \\subsection{Relations} \\label{relations}  The value of $L_{dust}/M_{gas}$ can be related to a depletion time for the dense gas and is often described as a star formation ``efficiency\". Values from the radiative transfer models range from 1.9 to 283, with a mean of 105 and a median of 78. A similar calculation for dense clumps in the Milky Way, as defined by HCN maps, yields similar values: in a range of 3.3 to 398, the mean value is $117\\pm20$, with a median of 97. This similarity indicates that modeling extreme starbursts as an ensemble of massive clumps similar to those in the Milky Way leads to a consistent result.  In the absence of detailed models, how should one interpret the dust temperature derived from fitting the SED? There is a strong correlation between the fitted dust temperature and the luminosity to mass ratio for the mean clump derived from the radiative transport models (Figure \\ref{tdvslm}). A least-squares fit in log space to the points with uncertainties in both axes yields the following: \\begin{equation} log (L_{dust}/M_{gas}) = (-4.55\\pm1.54) + (4.11\\pm0.97) log (\\td) \\end{equation} We have left SMM J02396 out of the fit, as it may represent a galaxy dominated by cirrus emission. It is plotted in Figure \\ref{tdvslm} and clearly deviates from the fit to the other sources. The value of $4.11\\pm0.97$ can be understood in terms of the Stefan-Boltzmann law, though we caution that $\\td$ is, at best, an average over a large range of real temperatures. A similar relation is found for massive, dense clumps in the Milky Way \\citep{wu09}.  Using the relation from \\citet{kennicuttrev} between star formation rate and far-infrared luminosity, we can compute a depletion time for the dense gas from  \\begin{equation} t_{dep}(y) = 5.6\\ee{9}  (L_{dust}/M_{gas})^{-1} \\\\ \\approx 2.0\\ee{14} \\td^{-4.11} \\end{equation} which ranges from 3.0\\ee9 yr for SMM J02396 down to 2.0\\ee7 yr for SDSS0338. The mean value is 4.7\\ee8 yr, or 1.1\\ee8 yr, excluding the value for SMM J02396.  The bottom plot in Figure \\ref{tdvslm} shows the log of the optical depth at 100 \\micron\\ from a single core versus the value of dust temperature from the simple fit to the SED. There is a strong anti-correlation. (We have again left SMM J02396 out of the fit, but it lies on the relation in this case.) Galaxies with SEDs suggesting warmer dust can be modeled with less opaque clumps as more luminosity is derived from a smaller amount of mass.  \\subsection{Comparison of Dust and CO} \\label{compare}  The dust masses we have calculated from the simple, optically thin isothermal fit (Table 4) assume an opacity given by models of dust in dense clouds in the Milky Way \\citep{ossen94}. The gas masses emerging from the more detailed models and given in Table 5 make the same assumption about opacity and also assume a gas to dust ratio of 100, as is usually assumed in the Milky Way. Both these assumptions could in principle be quite wrong for high-z EMGs. Masses computed from CO luminosities depend on a conversion to molecular hydrogen based on models of local ULIRGS, which is about a factor of 6 times lower than for galaxies like the Milky Way \\citep{annrev05}. Thus substantial uncertainties attend these mass estimates and consistency checks are worthwhile.  We have compared the masses derived from the simple fit to those from the detailed radiative transport models by using a gas to dust ratio of 100 to convert the former into gas masses. The mean ratio of simple estimate to model mass is 0.64 with surprisingly small dispersion. Thus, the simple fit gives a good estimate of the mass, even when the assumptions are invalid. As long as there is a data point at wavelengths long enough to be optically thin and in the Rayleigh-Jeans limit, a simple fit provides an estimate that can agree with those from more detailed models. If the models are correct, the masses from simple fits should be scaled up by a mere 1.56.  Taking the gas mass estimates based on CO luminosity from \\citet{annrev05}, we find that the mean ratio of the mass from CO to that from the models is 0.60, with a bit more dispersion (minimum is 0.15 for SMM J02396 and maximum is 1.50 for SMM J14011). The agreement between these two estimates is even more surprising, as there are many reasons that they could differ. The conversion from CO luminosity to mass is uncertain and depends on internal cloud conditions like density and temperature (e.g., \\citealt{dickman86}, \\citealt{solomon87}). The dust opacity model could be different or the gas to dust ratio could differ. For example, the two mass estimates could be reconciled if the gas to dust ratio is 167, which is probably within the range of uncertainty, even for the Milky Way (see \\citet{draine03} for discussion of abundance issues in grain models). The general consistency despite all these possible differences suggests that neither the standard conversion from CO luminosity to gas mass nor the gas to dust ratio is likely to be far off. The less satisfying possibility, that both conversions are wrong and the errors cancel out, cannot be ruled out.  Considering the possible differences in radiation environment, metallicity, etc., the consistency between the masses from CO and from dust emission is quite remarkable. We have used the submillimeter opacity values appropriate for coagulated, icy grains in dense molecular cores. The submillimeter opacity of the dust grains that characterize the diffuse interstellar medium in the Milky Way is less by a factor of 4.8, which would result in masses from dust emission that are, on average, larger than those inferred from CO by a factor of 3. Assuming that the CO conversion factor is exactly correct, the dust properties in EMGs would appear to lie between those of the diffuse ISM in the Milky Way and those of dense cores, but closer to the latter.  The remarkable thing about this comparison is that the masses from the models include only quite dense gas (typically $n \\geq \\eten{4}$ \\cmv, see Fig. 4), while the CO emission could easily arise in gas of much lower density. The agreement suggests that essentially all the molecular gas in these EMGs is in dense clumps.   \\subsection{Comparison to Other Models}  \\citet{narayanan09a} have recently proposed a model for the formation of high redshift submillimeter galaxies as the result of a starbursts triggered by major mergers in massive halos. They modeled the time evolution of such mergers and found that the most luminous observed submillimeter galaxies could be modeled as 1:1 mergers in the most massive dark halos. As the two galaxies spiral in toward the merger, SFRs of 100-200 \\msunyr\\ can be triggered. During final coalescence, they predicted a brief burst (5\\ee{7} yr) of very high SFR ($\\sim 2000$ \\msunyr). These predictions are broadly consistent with our numbers for SFR (Table 4) and our mean depletion time of (1.1-4.7)\\ee{8} yr. To make the most luminous submillimeter galaxies, they need to assume that all stellar clusters with ages less than 10 Myr are embedded in the material that is forming them, consistent with the assumptions of our modeling.  \\citet{narayanan09b} have discussed the interpretation of CO lines from such galaxies. They warn that the mass estimates from higher-J CO lines may underestimate the gas mass by a factor up to 2. This would be consistent with our result above, but we believe that the sources of uncertainty discussed in \\S 7.2 are equally important.  It is also interesting to compare our results to predictions of quite different kinds of models. \\citet{err03} modeled two of our sources with their model of cirrus dominated far-infrared emission. In SMM J02399, our data are consistent with their model, but in SMM J14011, their model prediction is about a factor of 10 too low. This result is generally consistent with the \\td\\ values in Table 5: SMM J02399 has $\\td = 29$ K, within the limit of 30 K that the dust achieves in their models; SMM J14011 has $\\td = 42$ K. On this basis, about half the sources in our sample could be consistent, with SMM J02396 as a particularly good case. Of course, models in which {\\it some} of the radiation arises in cirrus and some in dust surrounding the forming stars could be relevant. Since the cirrus model is the logical extreme of model 5 of RXJ0911.4, where we increased the ISRF, further constraints on the SED could help to constrain such models.  While our data do not bear directly on models using radiation from an AGN to heat a dusty torus, the general consistency of our models with dust emission from star formation bolsters the case for the embedded starburst models."
    },
    "0910/0910.2671_arXiv.txt": {
        "abstract": "A very interesting puzzle about the origin of electron and positron cosmic rays is deduced from the latests experimental results. We model the propagation of such cosmic rays in terms of a successfully tested two--zone propagation model. Several theoretical uncertainties -- like ones related to propagation -- are considered to study different types of electron and positron sources: dark matter annihilation, secondary production, and supernova remnants.\\\\ ",
        "introduction": "\\label{sec1}  During last decades, many efforts have been done to understand the nature of cosmic rays. Recently, the experimental results for the positron and electron signals obtained by PAMELA~\\cite{2009Natur.458..607A}, FERMI~\\cite{2009PhRvL.102r1101A}, and HESS~\\cite{2008PhRvL.101z1104A} have presented and confirmed a very interesting puzzle related to the origin of such cosmic rays.\\\\ The positron fraction, the ratio among number of positrons and number of electrons plus positrons, presents an increment for energies above $\\sim$50~GeV. Similar phenomenon appears in the total flux of electron and positron, which presents a change of behavior. For energies below 100~GeV, the total flux is similar to a power--law with a power index of $\\sim 3.3$. Above 100~GeV this limit, it changes to a power index of $\\sim 3$.\\\\ In this proceeding, we discuss and study the electron and positron propagation model, different cosmic rays sources as dark matter annihilation and astrophysical ones. We focus to study the supernova remnants as a feasible source of electrons.\\\\ ",
        "conclusions": "\\label{sec4} The improvement in current electron and positron measurements and data has revealed a very interesting puzzle. Many solution have been proposed during the years. We present the importance of theoretical uncertainties related to propagation and production of electron and positron cosmic rays~\\cite{2008arXiv0812.4272L, 2009A&A...501..821D,2008PhRvD..77f3527D}. We stress the necessity to re-estimate secondary and primary component, and to consider already known sources in order to discard/confirm possible presence of an undiscovered component, like the case of dark matter annihilations.  \\ack \\begin{small} Work supported by research grants funded jointly by Ministero dell'Istruzione, dell'Universit\\`a e della Ricerca (MIUR), by Universit\\`a di Torino (UniTO), by Istituto Nazionale di Fisica Nucleare (INFN) within the Astroparticle Physics Project and by the Italian Space Agency (ASI) under contract N I/088/06/0. R.L. acknowledges to T. Delahaye, F. Donato, N. Fornengo, J. Lavalle, P. Salati, and R. Taillet for unvaluable comments and suggestions. \\end{small}"
    },
    "0910/0910.1590_arXiv.txt": {
        "abstract": "A survey of Type~II supernovae explosion models has been carried out to determine how their light curves and spectra vary with their mass, metallicity and explosion energy.  The presupernova models are taken from a recent survey of massive stellar evolution at solar metallicity supplemented by new calculations at sub-solar metallicity. Explosions are simulated by the motion of a piston near the edge of the iron core and the resulting light curves and spectra are calculated using full multi-wavelength radiation transport. Formulae are developed that describe approximately how the model observables (light curve luminosity and duration) scale with the progenitor mass, explosion energy, and radioactive nucleosynthesis. Comparison with observational data shows that the explosion energy of typical supernovae (as measured by kinetic energy at infinity) varies by nearly an order of magnitude -- from $0.5$ to $4.0~\\times 10^{51}$~ergs, with a typical value of $\\sim 0.9 \\times 10^{51}$~ergs. Despite the large variation, the models exhibit a tight relationship between luminosity and expansion velocity, similar to that previously employed empirically to make SNe~IIP standardized candles.  This relation is explained by the simple behavior of hydrogen recombination in the supernova envelope, but we find a sensitivity to progenitor metallicity and mass that could lead to systematic errors. Additional correlations between light curve luminosity, duration, and color might enable the use of SNe~IIP to obtain distances accurate to $\\sim$20\\% using only photometric data. ",
        "introduction": "Type II supernovae (SNe~II) result from the explosion of massive stars that have retained their hydrogen envelope until their cores collapse to neutron stars or black holes. The most common events, the Type~II plateau supernovae, have a distinctive light curve, maintaining a nearly constant luminosity for $\\sim 100$~days, then suddenly dropping off.  Upcoming synoptic surveys should discover millions of these events out to redshifts of a few.  These observations will probe massive stellar evolution in a broad range of galactic environments, and may also be used to measure cosmological distances and test host galaxy dust properties.  SNe~II are diverse transients in terms of their luminosities, durations, and expansion speeds. By modeling the observed light curves and spectra, one can constrain the physical properties of the explosion such as the ejected mass, explosion energy, and presupernova radius. In a pioneering study, \\cite{Hamuy_SNIIP} analyzed 16 observed supernovae and derived masses between $15$ and $50~\\Msun$ and presupernova radii ranging from $70$ to $600~$R$_\\odot$.  Unfortunately these inferred values seem implausible, especially the high masses.  Direct observations of the progenitors of nearby SNe~II from pre-explosion images indicate masses of only $8-15~\\Msun$ \\citep{Smartt_prog}, while modern stellar evolution models predict a mass range of $\\sim 12-25$~\\Msun\\ \\citep[e.g.,][]{Woo02,Woo07}.  The overly large values inferred by \\cite{Hamuy_SNIIP} likely reflect deficiencies in the theoretical models used in the light curve analysis \\citep{LN_85}. More recent models tailored to individual events have returned more reasonable numbers \\citep[e.g.][]{Utrobin_05cs, Baklanov}.  Observations of SNe~II are also  used to measure cosmological distances.  In terms of brightness alone, they are poor standard candles with luminosities varying by more than an order of magnitude, but various methods can be used to standardized them.  Most previous studies were of the Baade-Wessellink type, e.g., the expanding photosphere method \\citep{Kirshner_EPM,Eastman_EPM,Dessart_EPM,Dessart_05cs,Jones_EPM} and the spectral-fitting expanding atmosphere method \\citep[SEAM,][]{Mitchell_1987A,Baron_99em}. Both approaches require detailed atmospheric modeling. More recently, \\cite{Hamuy_SCM} suggested a much simpler approach using an empirical correlation between plateau luminosity and expansion velocity (as measured from the Doppler shift of spectral lines).  This standardized candle method has been extended out to moderate redshifts with promising results \\citep{Nugent_SCM, Dovi_SCM}.  However, the physical underpinnings of the  method and its potential vulnerability to systematic error have not yet been fully explored.  Here we present a new survey of SNe~IIP models based on the one-dimensional explosion of realistic progenitor star models of various masses, metallicities, and explosion energies.  We calculate the broadband light curves and the detailed spectral evolution using a code that includes a full multi-wavelength solution to the radiation transport problem. The models allow us to explore the light curves dependence on model parameters, which we compare to analytic expectations.  We also reproduce and illuminate the physics leading to the standardized candle relation.  \\clearpage ",
        "conclusions": "\\begin{figure} \\includegraphics[width=3.3in]{f17.eps} \\caption{Relationship between the plateau duration (assuming zero \\Nifs) and the luminosity at day 50.  Circles denote solar metallicity models, squares $0.1$~solar models, and the size of the symbol is proportional to the progenitor star mass. \\label{Fig:ltp} } \\end{figure}  We explored the light curves and spectra of SNe~II models with various progenitor masses, metallicities, and explosion energies.  We found that explosions with energies $0.3-4.8$~B of stars with initial masses in the range $12-25~\\Msun$ can explain the observed range of luminosities, velocities, and light curve durations of most SNe~IIP. For existing and future observational surveys, the model results should be useful for inferring the progenitor star properties, explosion energies, distances, and dust extinction of observed events.  This study, as have previous studies, quantified how the basic supernova parameters (\\Mej, $R_0$, and $E$) affect the light curves. We also highlighted the important role of two additional parameters: the radioactive \\Nifs\\ mass and the envelope helium abundance. The presence of \\Nifs\\ extends the plateau duration, but typically does not affect the luminosity on the plateau at times $t \\la 50$~days. The neglect of the effect of \\Nifs\\ may be the main reason why \\cite{Hamuy_SNIIP}, in his analysis of 16 SNe~IIP, inferred implausibly large ejecta masses (up to $50~\\Msun$).  In that study, the longer plateau duration would have to be accounted for by an increased diffusion time, and hence larger ejecta mass.  Here we presented analytical formulae which may be useful in accounting for the effects of \\Nifs\\ on the plateau.  The models confirm the standard candle method of calibrating SNe~IIP and illuminate its physical origin.  The method is a promising way to determine distances to SNe~IIP, with a clear physical explanation in terms of the ionization physics of hydrogen. On the other hand, the models raise some concerns about systematic errors.  Progenitors with different masses or metallicities lie on differently normalized relations in Figure~\\ref{Fig:Hamuy}.  If the progenitor population at high redshift  has different demographics than that at low redshift (as might be expected) a systematic bias may be introduced into distance measurements.  The effect of going from $Z=1$ to $Z=0.1$ solar metallicity in the models is at the 0.1~mag level. It may be possible to reduce these errors by using color information from the light curve.  We find a correlation between plateau luminosity and plateau duration which could be useful in roughly calibrating SNe~IIP luminosities using only photometric data (to about 20\\% in distance).  This correlation reflects the fact that in the models one parameter, the explosion energy, primarily controls both the light curve brightness and duration, while the progenitor star properties play a secondary role.  The validity of such a relation needs to be empirically  checked, as the scatter will be smaller or larger depending on whether the bulk of SNe~IIP arise from a narrower or wider range of progenitor masses and radii than that considered here.  In practice, the relation also  needs to take into account the effect of \\Nifs\\ on extending the plateau duration.  Correction for dust extinction remains a difficult issue for determining the distances to SNe~IIP.  The models provide some theoretical guidance as to the intrinsic color evolution of SNe-IIP light curves, however their accuracy may be limited by the assumptions in the radiative transfer, and are sensitive to variations in the envelope metallicity. On the other hand, one could try to invert the problem.  Assuming the cosmological parameter are accurately constrained by other means, one could use the standard candle method to solve for the dust extinction of SNe~IIP, thus providing an estimate of the variation of dust properties with galactic environment and redshift."
    },
    "0910/0910.1842_arXiv.txt": {
        "abstract": "The Pierre Auger Observatory's (PAO) shower profile measurements can be used to constrain the chemical composition of the ultra-high energy cosmic ray (UHECR) spectrum. In particular, the PAO's measurements of the average depth of shower maximum and the fluctuations of the depth of shower maximum indicate that the cosmic ray spectrum is dominated by a fairly narrow distribution (in charge) of heavy or intermediate mass nuclei at the highest measured energies ($E \\gsim 10^{19}$ eV), and contains mostly lighter nuclei or protons at lower energies ($E \\sim 10^{18}$ eV). In this article, we study the propagation of UHECR nuclei with the goal of using these measurements, along with those of the shape of the spectrum, to constrain the chemical composition of the particles accelerated by the sources of the UHECRs. We find that with modest intergalactic magnetic fields, 0.3~nG in strength with 1~Mpc coherent lengths, good fits to the combined PAO data can be found for the case in which the sources accelerate primarily intermediate mass nuclei (such as nitrogen or silicon). Without intergalactic magnetic fields, we do not find any composition scenarios that can accommodate the PAO data. For a spectrum dominated by heavy or intermediate mass nuclei, the Galactic (and intergalactic) magnetic fields are expected to erase any significant angular correlation between the sources and arrival directions of UHECRs. ",
        "introduction": "The chemical composition of the ultra-high energy cosmic ray (UHECR) spectrum has long been a topic of great interest~\\cite{Stecker:1969fw,Puget:1976nz,Stecker:1998ib,us,debate,magnetic}. Until recently, however, very little was known about the nature of these particles. On one side of the debate, the so called Hillas criterion~\\cite{hillas} gives a preference for the electromagnetic fields of cosmic ray sources to accelerate heavy nuclei to higher energies than protons or light nuclei. On the other side, it has been argued that the angular correlations reported by the Pierre Auger Observatory (PAO)~\\cite{anisotropy}, as well as features in the shape of the UHECR spectrum~\\cite{berezinsky}, suggest that these particles consist largely of protons. None of these arguments, however, has yet settled the question of what types of particles make up the UHECR spectrum.  Data from the Pierre Auger Observatory (PAO), however, is offering increasingly powerful insights into this question. Firstly, the spectral shape predicted for the UHECR all-particle spectrum depends not only on the injected spectrum and spatial distribution of the sources, but also on the chemical composition that is injected from the sources of the highest energy cosmic rays. As the PAO measures the UHECR spectrum with increasing precision~\\cite{spectrum}, this information can be used to constrain the chemical composition of these particles~\\cite{spectrumcomposition}. Furthermore, the PAO is capable of performing several measurements that can be used to directly or indirectly determine the chemical composition of UHECRs as they enter the Earth's atmosphere. Among these empirical tools are the measurements of the average depth of shower maximum, $\\langle X_{\\rm max} \\rangle$, and the RMS variation of this quantity, $\\mathrm{RMS}(X_{\\rm max})$. On average, proton-induced showers reach their maximum development, $\\langle X_{\\rm max} \\rangle$, deeper in the Earth's atmosphere than do showers of the same energy generated by heavier nuclei. Accompanying this result, the shower to shower fluctuation of $X_{\\rm max}$ about the mean, $\\mathrm{RMS}(X_{\\rm max})$, is larger for proton-induced showers than for iron-induced showers of the same energy. As a result, measurements of both $\\langle X_{\\rm max} \\rangle$ and $\\mathrm{RMS}(X_{\\rm max})$ can be used to infer the average chemical composition of the UHECRs as a function energy.  Very recently, the PAO collaboration has announced their first measurements of $\\mathrm{RMS}(X_{\\rm max})$~\\cite{xmax,ungerSocor}. These measurements, along with those of $\\langle X_{\\rm max} \\rangle$, imply that the UHECR spectrum contains a large fraction of heavy or intermediate mass nuclei, especially at the highest energies measured. Furthermore, the small values of $\\mathrm{RMS}(X_{\\rm max})$ measured by the PAO also imply that the composition of the UHECR spectrum is relatively narrowly distributed at the highest measured energies, containing little or no protons or light nuclei. In this way, the new $\\mathrm{RMS}(X_{\\rm max})$ measurements not only confirm and reinforce the conclusions drawn from earlier average depth of shower maximum measurements, but also provide complementary information that enables one to constrain the distribution of the various chemical species present within the UHECR spectrum.  The remainder of this article is structured as follows. In Sec.~\\ref{prop} we review the physics of UHECR nuclei propagation. In Sec.~\\ref{results}, we calculate the spectrum and chemical composition of the UHECR spectrum for various choices of the injected chemical species and compare these results to measurements of the PAO, neglecting the effects of intergalactic magnetic fields. In Sec.~\\ref{resultsb}, we include magnetic fields in our calculation, and show that fields with stengths and coherent scales $\\sim$($B$/0.3~nG)$\\times$($L_{\\rm coh}$/1~Mpc)$^{1/2}$, are strongly preferred to accomodate the PAO's data. The best fits found are for cases in which the sources of the UHECRs inject mostly intermediate mass nuclei, such as nitrogen or silicon. Finally, in Sec.~\\ref{conclusions}, we discuss the implications of our results and summarize our conclusions. ",
        "conclusions": "\\label{conclusions}  In this article, we have studied the propagation of ultra-high energy cosmic ray nuclei in an attempt to find scenarios that provide a good fit to the observations of the Pierre Auger Observatory (PAO), including the measured spectrum, average depth of shower maximum ($\\langle X_{\\rm max} \\rangle$), and the RMS fluctuations of the depth of shower maximum ($\\mathrm{RMS}(X_{\\rm max})$).  We have found that for all values of injected maximum energy and spectral index considered for the UHECR sources, the PAO's data, for energies above 10$^{18.5}$~eV, is best accommodated in cases in which a homogeneous distribution of sources inject largely intermediate mass nuclei, ranging from roughly nitrogen to silicon. This result is found for all hadronic models considered, with the best fit results being found using the EPOS model. Furthermore, within this same range of source parameters, the presence of $\\sim$($B$/0.3~nG)$\\times$($L_{\\rm coh}$/1~Mpc)$^{1/2}$ intergalactic magnetic fields lead to far better fits than are found in the absence of cosmic ray deflection. No good fits to the PAO's measurements were found without significant intergalactic magnetic fields.\\footnote{Although magnetic fields in the vicinity of the sources could also plausibly account for this conclusion.} Along with these results we note that for the EPOS model, whose results are responsible for our best fits, no good fits were found with the in any scenario considered in which the sources of the UHECRs inject a significant ($\\gsim 10\\%$) fraction of protons. Should the PAO data points all be shifted up in energy by approximately 20\\%, as has been suggested \\cite{Engel:2007cm}, the best fit results, for all cases quoted here, are found to shift towards a heavier nuclei composition.  These results have a number of interesting and important consequences. Firstly, if the cosmic ray spectrum consists largely of intermediate mass nuclei at the highest energies, as our results suggest is the case, one would not expect any significant correlation between the arrival directions and sources of such particles. For a $\\mu$G-scale Galactic Magnetic Field with a correlation lenth of $\\sim$1 kpc, a $6\\times 10^{19}$ eV nucleus with $Z \\approx 10$ would be deflected by on the order of 10$^{\\circ}$, in addition to any deflection resulting from intergalactic fields. This conclusion, perhaps, suggesting difficulties for the prospects of UHECR astronomy, for all but the most bright local sources.  Under the assumption that the UHECRs consist uniquely of protons, a ``guaranteed'' flux of UHE neutrinos is predicted to be produced as a result of their propagation. This signal, known as the cosmogenic neutrino flux~\\cite{cosmogenic}, is potentially within the reach of experiments such as IceCube~\\cite{icecube}, ANITA~\\cite{anita}, and the PAO itself~\\cite{paoneutrino}. If the the UHECR spectrum consists of heavy nuclei, however, the cosmogenic flux can be considerably suppressed~\\cite{cosmogenicnuclei}. In the scenarios that we have found to fit the PAO's data, the cosmogenic neutrino spectrum is predicted to be suppressed by an order of magnitude or more relative to the all-proton case~\\cite{suppress}. The neutrino flux from the sources of the UHECRs themselves is also expected to be suppressed in these scenarios~\\cite{Anchordoqui:2007tn}.  The authors note that since the submission of this work, new results on the UHECR composition from $\\langle X_{\\rm max} \\rangle$ measurements have been published from the HiRes experiment \\cite{Abbasi:2009nf}. These results appear at first sight to sit in contradiction to those from the PAO. Along with the apparent contradiction between these data sets, other recent results from the Yakutsk experiment fail to resolve this confusion, with different analysis of the data arriving at different results for the UHECR composition \\cite{Glushkov:2007gd,Knurenko:2007yv}, an issue whose origin may lie in the different methods used to asign energy to the events \\cite{private}. Due to the strength of the conclusions that have been possible from PAO results, the conflict between these sets of results are considered a pressing issue which we hope the experimental community will thoroughly investigate.  \\vspace{1.0cm}  {\\it Acknowledgements:} We would like to thank both Subir Sarkar and Felix Aharonian for their invaluable contributions to this work. We would also like to thank Michael Unger for useful discussions. DH is supported by the US Department of Energy, including grant DE-FG02-95ER40896, and by NASA grant NAG5-10842."
    },
    "0910/0910.1629_arXiv.txt": {
        "abstract": "We present the observations of GRB090510 performed by the {\\it Fermi} Gamma-Ray Space Telescope and the {\\it Swift} observatory. This is a bright, short burst that shows an extended emission detected in the GeV range. Furthermore, its optical emission initially rises, a feature so far observed only in long bursts, while the X-ray flux shows an initial shallow decrease, followed by a steeper decay. This exceptional behavior enables us to investigate the physical properties of the GRB outflow, poorly known in short bursts. We discuss internal shock and external shock models for the broadband energy emission of this object. ",
        "introduction": "With the availability of a relatively large sample of Gamma-Ray Bursts (GRBs), we came to recognize that they comprise of two large classes \\citep{kou93}: the so-called short-hard GRBs (duration $ \\lesssim 2$~s) and the long-soft ones ($ \\gtrsim 2$~s). There is now increasing consensus that the observed dichotomy  among long/short GRBs may indicate diverse initial physical conditions and progenitors. Long GRBs are associated with the demise of massive stars \\citep{fer06}. Instead, short GRBs often occur in early-type galaxies \\citep{zha09,geh09,fon09}. This supports their interpretation in terms of compact object mergers.   Crucial information on GRBs is revealed by their afterglows, which can be monitored by {\\it Swift} \\citep{geh04} in the optical and the X-ray range as soon as $\\sim100$~s after the burst. In this paper, we present the study of the short GRB090510 with {\\it Swift} and {\\it Fermi} in a broad energy range, which extends from the optical up to a few GeV.    We report our observations and analysis in \\S 2; in \\S 3 we propose two different interpretations, in \\S 4 we draw our conclusions. Hereafter, we use the conventions $X = 10^n X_n$ for cgs units and $F\\propto t^{-\\alpha} \\nu^{-\\beta}$, where $F$ is energy flux, $t$ is the time from the trigger of ${\\it Swift}$ Burst Alert Telescope \\citep{bar05a} and $\\nu$ the frequency. Errors are reported at $1\\sigma$, unless otherwise specified. We assume a cosmology in which $H_0=71$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_m = 0.27$, $\\Omega_\\lambda=0.73$. All the fluxes, times and frequencies are measured in the observer's frame. ",
        "conclusions": "We have reported the {\\it Swift} and {\\it Fermi} observations of the short GRB090510, an event endowed with bright prompt and afterglow emission, and detected in the GeV range up to 200~s after the trigger. The initial X-ray emission shows a slow decay up to $\\simeq 1.4$~ks, after which it quickly drops. The optical flux peaks at $\\simeq 1.6$~ks. We have explored two scenarios to explain the observed behaviors.\\\\ In the first scenario, the early flux detected by XRT and LAT is due to IS, while the optical rise is the onset of FS emission or the transit of $\\nu_m$. This interpretation does not require extremely high values of Lorentz factor, should the density of the environment be very low. We also find that reasonable values for the physical parameters can lead to the observed properties, which might favour the model, although some fine tuning is necessary. The second scenario assumes that the FS produces the full spectrum of the emission, observed from the optical to the GeV band. The $\\gamma$-ray, X-ray and optical spectrum can be reproduced by the template FS spectral model and the required physical parameters are plausible. Although the simple FS model fails to reproduce the observed temporal behavior, extensions of this model could accommodate the temporal mismatch. In order to identify the origin of the GeV component of GRBs like 090510, more case studies will be necessary. Fortunately, we have very promising prospects for other simultaneous {\\it Fermi} and {\\it Swift} observations of short GRBs, which will provide us with more measurements to shed light on the properties of this class of events."
    },
    "0910/0910.5520_arXiv.txt": {
        "abstract": "The first results from observations of the high mass X-ray binary LS 5039 using the \\fermi\\ Gamma-ray Space Telescope data between 2008 August and 2009 June are presented.  Our results indicate variability that is consistent with the binary period, with the emission being modulated with a period of $3.903 \\pm 0.005$ days; the first detection of this modulation at GeV energies.  The light curve is characterized by a broad peak around superior conjunction in agreement with inverse Compton scattering models.  The spectrum is represented by a power law with an exponential cutoff, yielding an overall flux (100 MeV -- 300 GeV) of 4.9 $\\pm$ 0.5(stat) $\\pm$ 1.8(syst) $\\times$10$^{-7}$ photon cm$^{-2}$ s$^{-1}$, with a cutoff at 2.1 $\\pm$ 0.3(stat) $\\pm$ 1.1(syst) GeV and photon index $\\Gamma$ = 1.9 $\\pm$ 0.1(stat) $\\pm$ 0.3(syst).  The spectrum is observed to vary with orbital phase, specifically between inferior and superior conjunction.  We suggest that the presence of a cutoff in the spectrum may be indicative of magnetospheric emission similar to the emission seen in many pulsars by {\\em Fermi}. ",
        "introduction": "LS 5039, PSR B1259$-$63 and \\lsi\\ are the only binaries, with high-mass companions, long known to be spatially coincident with sources detected at energies greater than 100 MeV, e.g., those listed in the Third Energetic Gamma-Ray Experiment (EGRET) catalog \\citep{1999ApJS..123...79H}. The latter binary was detected by the Large Area Telescope (LAT) on the \\fermi\\ mission, confirming it as a GeV gamma-ray source and finding variability that is consistent with modulation on the binary period of 26.6 $\\pm$ 0.5 days \\citep{2009ApJ...701L.123A}. This constituted the first detection of orbital periodicity in high-energy (HE) gamma rays (20 MeV -- 100 GeV). In this Letter, we present the results of \\fermi\\ observations of LS 5039.  LS 5039 is one of a handful of X-ray binaries that have been detected recently at very HE $\\gamma$-rays; results from $\\sim70$ hr of observations distributed over many orbital cycles have been presented by \\citet{Aharonian:2005nj, Aharonian:2006qw}.  These observations yielded a modulation of the very high energy (VHE, >100 GeV) gamma-ray flux with a period of 3.9078$\\pm$0.0015 days \\citep{Aharonian:2006qw}, consistent with the orbital period as determined by \\citet{Casares:2005gg} from optical spectroscopy. Short timescale variability displayed on top of this periodic behavior, both in flux and spectrum, was also reported.   The nature of the LS 5039 compact object is unknown: a black hole and a neutron star have been invoked as possible compact object companions in a slightly eccentric ($e\\sim0.35$), 3.90603$\\pm$0.00017 day orbit around the O6.5V star \\citep{Casares:2005gg}. The discovery of a jet-like radio structure in LS 5039 coincident with an X-ray and EGRET source prompted a microquasar interpretation \\citep{2000Sci...288.2340P}. Variability in the EGRET source could not be established precisely \\citep{2001A&A...370..468T, 2003ApJ...597..615N}. Recently it has been shown that the X-ray flux is modulated on the orbital period \\citep{Takahashi:2008vu}.  \\citet{2008A&A...481...17R} provided Very Long Baseline Array (VLBA) radio observations of LS 5039 with morphological and astrometric information at milliarcsecond scales that cannot easily be explained by a microquasar scenario. \\citet{2005A&A...430..245M} assessed the low X-ray state, showing the absence of accretion features in the X-ray spectra.  Thus, measurements at radio and VHE $\\gamma$-rays in the cases of \\lsi\\ \\citep{2006smqw.confE..52D, 2008ApJ...684.1351A} or PSR B1259$-$63 \\citep{2005A&A...442....1A}, whose overall spectral energy distribution is similar to that of LS 5039, gave the impression that all three systems are different realizations of the same scenario: a pulsar-massive star binary \\citep[][]{Dubus:2006lc}.  Theoretical computations of the gamma-ray emission in both compact object scenarios have been made, with gamma-rays produced by inverse Compton (IC) scattering of the stellar light by VHE electrons accelerated in the vicinity of the compact object. In the case of the black hole companion, HE and VHE emission would result from particles accelerated in the jet \\citep[][]{2007A&A...464..259B, 2007APh....27..278B, 2008MNRAS.383..467K}. Alternatively, it would involve the relativistic wind of a young, rotation-powered pulsar, either as a result of particle acceleration in the wind interaction region \\citep{Dubus:2006lc} or by processes within the pulsar wind \\citep{2008ApJ...674L..89S, 2008APh....30..239S, 2008A&A...488...37C}.  \\begin{figure*} \\centering \\includegraphics[width=0.9\\linewidth,angle=0,clip]{f1.eps} \\caption{Left: the smoothed counts map for 100~MeV--300~GeV of a 10\\degr $\\times$10\\degr\\ region around the LS 5039 location (marked with ``$\\times$''); the dashed black circle indicates the 0.$^{\\circ}$925 timing analysis aperture.  A 0.$^\\circ$3 gaussian smoothing function was applied to the 0.$^\\circ$1 bins. PSR J1826$-$1256 is marked with ``+''. The exposure varies by less than 7\\% across the field at a representative energy of 10 GeV. Right: residuals of left panel after subtracting all modelled sources excluding LS 5039 and excluding events which arrive during the peaks in the pulsar phase cycle of PSR J1826$-$1256 (see section~\\ref{sec:spectra}). The white circles indicate the 68\\% LAT containment region at three energies: 0.3, 1, and 10 GeV.  Note that the color scales are different for the two panels.} \\label{CountsMap} \\end{figure*} ",
        "conclusions": "The initial association of LS 5039 with the EGRET source 3EG J1824$-$1514 \\citep{Paredes:2000ci} had remained tentative due to the large EGRET error circle and the lack of timing signatures. The association was bolstered by the discovery of point-like, modulated gamma-ray emission above 100 GeV \\citep{Aharonian:2005nj,Aharonian:2006qw}. The {\\em Fermi} observations enabled the detection of an orbital modulation, indicating that the binary is also a source of gamma rays above 100 MeV. The periods determined independently from the {\\em Fermi}-LAT and \\hess\\ data are self consistent and compatible with the binary period obtained from radial velocity measurements of the companion \\citep{2009ApJ...698..514A}. This is only the second high mass X-ray binary after \\lsi\\ \\citep{2009ApJ...701L.123A} to become a confirmed emitter in the HE gamma-ray domain.  The short orbital period means that the compact object passes within a stellar radius of the surface of the $kT_\\star \\approx 3$ eV, $R_\\star \\approx 10 R_\\odot$ companion \\citep{Casares:2005gg}. Gamma rays emitted in the vicinity of the compact object with energies above the 30 GeV threshold inevitably pair produce with stellar photons \\citep[see e.g.][]{1987ApJ...322..838P,1995SSRv...72..593M,Bottcher:2005zx,Dubus:2006lr}. Emission in the {\\em Fermi} range is largely unaffected by absorption and should allow easier identification of the intrinsic spectrum and variability of the source.  However, it can be affected by cascading of higher-energy photons.  The HE modulation peaks at $\\phi \\sim$0.0--0.1, close to superior conjunction ($\\phi=0.06$) while the VHE modulation peaks close to inferior conjunction ($\\phi=0.72$). The phase difference can be explained mostly as a result of the competition between IC scattering on HE electrons and VHE pair production, assuming the companion star to be close to the compact object. The star provides target photons for both processes with the radiation density varying by a factor 4 along the eccentric orbit. IC scattering will vary with radiation density but, because the source of seed photons is anisotropic, the flux will also depend on the geometry seen by the observer in non-trivial ways \\citep{2008MNRAS.383..467K, 2008APh....30..239S}. The emission is enhanced (reduced) when the highly relativistic electrons seen by the observer encounter the seed photons head-on (rear-on), i.e., at superior (inferior) conjunction. Inversely, VHE absorption due to pair production will be maximum (minimum) at superior (inferior) conjunction. The phases of minimum and maximum flux in {\\em Fermi} as well as the anti-correlation with \\hess\\ are consistent with these expectations, suggesting IC scattering is the dominant radiative process above 100 MeV with the additional effect of pair production above 30 GeV \\citep{Bednarek:2007qd,Dubus:2007oq,2008ApJ...674L..89S}. It is, however, yet unclear whether the IC VHE emission is mainly produced already in the pulsar wind zone of the system, given the high opacity already found therein for particles accelerated at the pulsar site \\citep{2008APh....30..239S} or as a result of particle acceleration at the shock formed in the wind collision region \\citep{Dubus:2006lr} or even well beyond the system \\citep{2008A&A...489L..21B}.  However, the extension of these pictures from the TeV into the GeV domain is undermined by the presence of an exponential cutoff at a few GeV in the {\\em Fermi} spectrum. A cutoff at $\\sim$6.3 GeV was also observed in \\lsi\\, indicating that this may be a common spectral feature in this class of source. The companion star in \\lsi\\ is a Be star with a dense equatorial disk. Passage of the compact object through this disk might have explained the exponential cutoff: for instance, this would crush a putative pulsar wind nebula closer to the neutron star, increasing synchrotron losses and introducing a strong dependence with orbital phase of the electron energy distribution \\citep{Dubus:2006lc}. But there is no such large, systematic contrast in the density of the stellar wind from the O6.5V star along the orbit in LS 5039. The presence of a similar cutoff in both systems argues against explanations related to the properties of the orbit or companion star.  The cutoff seems to require that the radiative process or radiating electrons be different between the HE and VHE domains, in disagreement with the picture presented above.   An intringuing possibility is that the emission in the {\\em Fermi} range from both LS 5039 and \\lsi\\ is magnetospheric emission as seen in the dozens of pulsars that have now been detected by {\\em Fermi}. The typical {\\em Fermi} pulsar emission has a hard power-law spectrum with a photon index $\\approx$ 1.5 and an exponential cutoff at $\\approx 2.5$ GeV. In this case, the cutoff energy is thought to be set by the balance between acceleration and losses to curvature radiation. The emission should in such a case be pulsed, although the orbital motion makes it exceedingly difficult to find in the {\\em Fermi} data with no prior knowledge of the spin period. Pulsar wind emission would dominate in the neighbouring hard X-ray and VHE bands, supported by their similar orbital modulations \\citep{Hoffmann:2008ys,Takahashi:2008vu}. It is to be noted that there is, however, an issue associated with having dominant magnetospheric emission in the HE band: magnetospheric emission is usually thought to be due to curvature radiation and has no obvious reason to be modulated with the orbital motion (although magnetospheric emission from pulsars in close binaries has hardly been modeled). The dense photon environment of the binaries may perhaps introduce differences compared to gamma-ray magnetospheric emission from isolated pulsars (for instance in the pair multiplicity). Hence, the spectrum suggests a magnetospheric emission interpretation, which is hard to reconcile with the IC scattering interpretation suggested by the modulation.  Future HE and VHE should lead to better constraints on the variability of the spectral parameters along the orbit.  This will help resolve whether there are several components and what their relative amplitudes are."
    },
    "0910/0910.0075_arXiv.txt": {
        "abstract": "This is a preliminary report on surface photometry of the major fraction of known globular clusters, to see which of them show the signs of a collapsed core. We also explore some diversionary mathematics and recreational tables. ",
        "introduction": "A focal problem today in the dynamics of globular clusters is core collapse. It has been predicted by theory for decades \\citep{hen61,lyn68,spi85}, but observation has been less alert to the phenomenon. For many years the central brightness peak in M15 \\citep{kin75,new78} seemed a unique anomaly. Then \\citet{aur82} suggested a central peak in \\object{NGC 6397}, and a limited photographic survey of ours \\citep[Paper I]{djo84} found three more cases, \\objectname{NGC 6624}, \\objectname[M 15]{NGC 7078}, and \\object[Cl 1938-341]{Terzan 8}), whose sharp center had often been remarked on \\citep{can78}.  As an example of how the new AASTeX object tagging macros work, we will cite some of the ``Superlative'' objects mentioned in section 10 of Trimble's (1992) review of astrophysics in the year 1991.  The youngest star yet found was \\object[\\[JCC87\\] IRAS 4]{IRAS 4} in \\objectname{NGC 1333}. \\object{70 Oph} was found to be the longest period spectroscopic binary. The most massive white dwarf was \\object{GD 50}, estimated at 1.2 solar masses. The first neutral hydrogen found in a globular cluster was \\object{NGC 2808} while the \\objectname[SDSS J093401.92+551427.9]{I Zw 18} retained the record for metal deficiency.  However, another low metallicitity galaxy was \\object{UGC 4483} in the \\objectname{M 83} group. The largest redshift \\object[PC 1247+3406]{source} in 1991 was found at z=4.897.  Lastly, what paper would be complete without a mention of the \\object[M1]{Crab nebula}! ",
        "conclusions": ""
    },
    "0910/0910.3645.txt": {
        "abstract": "We estimate the fraction of mass that is composed of compact objects in gravitational lens galaxies. This study is based on microlensing measurements (obtained from the literature) of a sample of 29 quasar image pairs seen through 20 lens galaxies. We determine the baseline for no microlensing magnification between two images from the ratios of emission line fluxes. Relative to this baseline, the ratio between the continua of the two images gives the difference in microlensing magnification. The histogram of observed microlensing events peaks close to no magnification and is concentrated below 0.6 magnitudes, although two events of high magnification, $\\Delta m \\sim 1.5$, are also present. We { study the likelihood of the microlensing measurements using} frequency distributions obtained from simulated microlensing magnification maps for different values of the fraction of mass in compact objects, $\\alpha$.  The { concentration of microlensing measurements close to} $\\Delta m \\sim 0$  can be {explained} only by simulations corresponding to very low values of $\\alpha$ (10\\% or less). {A maximum likelihood test yields $\\alpha=0.05_{-0.03}^{+0.09}$ }(90\\% confidence interval) {for a quasar continuum source of intrinsic size $r_{s_0}\\sim 2.6 \\cdot 10^{15} \\rm\\, cm$}. This estimate is valid in the $0.1 - 10M_\\odot$ range of microlens masses. {We study the dependence of the estimate of $\\alpha$ with $r_{s_0}$, and find that $\\alpha \\lesssim 0.1$ for $r_{s_0}\\lesssim 1.3 \\cdot 10^{16} \\rm\\, cm$. High values of $\\alpha$ are possible only for source sizes much larger than commonly expected ($r_{s_0}>>2.6 \\cdot 10^{16} \\rm\\, cm$).} Regarding the current controversy about Milky Way/LMC and M31 microlensing studies, our work supports the hypothesis of a very low content in MACHOS (Massive Compact Halo Objects). In fact, according to our study, quasar microlensing probably arises from the normal star populations of lens galaxies and there is no statistical evidence for MACHOS in the dark halos. ",
        "introduction": "The composition of matter in the halos of galaxies is a central problem in astrophysics. During the last 10 years, several observational projects have used gravitational microlensing \\citep{1986ApJ...304....1P} to probe the properties of the halos of the Milky Way (MACHO, \\citealp{2000ApJ...542..281A}; EROS, \\citealp{2007A&A...469..387T}) and M31 (POINT-AGAPE, \\citealp{2005A&A...443..911C}; MEGA, \\citealp{2006A&A...446..855D}). These experiments are based on the detection of magnification in the light-curve of a source induced by an isolated point-like (or binary) object passing near the observer's line of sight. From the successful detection of a number of microlensing events these collaborations have estimated the fraction of the halo mass that is composed of lensing objects, $\\alpha$. However, the reported results disagree. For the Milky Way's halo the measurements of the MACHO collaboration \\citep{2000ApJ...542..281A} correspond to a halo fraction of $0.08<\\alpha<0.50$ while EROS \\citep{2007A&A...469..387T} obtains $\\alpha<0.08$. On the other hand, re-analysis of publicly available MACHO light-curves \\citep{2004MNRAS.352..233B} leads to results similar to those reported by EROS (however, see also the counter-report by \\citealp{2005MNRAS.359..464G}). For M31 the AGAPE \\citep{2005A&A...443..911C} collaboration finds a halo fraction in the range $0.2<\\alpha<0.9$, while MEGA \\citep{2006A&A...446..855D} finds a limit of $\\alpha<0.3$.   The method applied to the Milky Way and M31 { can be extended} to the extragalactic domain by observing the microlensing induced by compact objects in the lens galaxy halo in images of multiply imaged quasars { (quasar microlensing; \\citealp{1979Natur.282..561C}, see also the review by \\citealp{2006glsw.book..453W}). Interpreting the light-curves of QSO 2237+0305, \\citet{1991AJ....102.1939W} suggest that the monitoring of microlensing variability can provide a measure of the optical depth in compact objects and in the smooth mass distribution. \\citet{1996MNRAS.283..225L} proposed a statistical approach to the determination of the mass density in compact objects based on the comparison between the observed and simulated magnification probability distributions. Microlensing can also be measured from a single-epoch snapshot of  the anomalous flux ratios induced by this effect between the images of a lensed quasar (\\citealp{1995ApJ...443...18W}; see also \\citealp{2002ApJ...580..685S}). \\citet{2004IAUS..220..103S} explore the practical application of this idea by using a sample of eleven systems with measured flux anomalies. Other quasar microlensing studies of interest for the present study are aimed at the determination of accretion disk sizes (e.g., the studies based in relatively large samples by \\citealp{2007ApJ...660..146P}, \\citealp{2007arXiv0707.0305M} and references therein).}  In practice, { the study of extragalactic microlensing} meets significant obstacles, in particular (e.g., \\citealp{2004ApJ...605...58K}) larger time-scales for microlensing variability and lack of a baseline for no magnification needed to detect and to quantify microlensing (see, however, the time variability based studies of several individual systems in \\citealp{2008ApJ...689..755M} and references therein). In addition, microlensing by an isolated object is not a valid approximation. Microlensing at high optical depth should be modeled (e.g., by simulating magnification maps; see \\citealp{1992grle.book.....S}).  We avoid these obstacles by setting the baseline of no microlensing magnification using the narrow emission lines (NELs) in the spectra of lensed quasar images (\\citealp{2004IAUS..220..103S} follow a similar approach but using theoretical models to define the baseline).  It is generally expected that the regions where NELs originate are very large (compared with the continuum source) and are not affected by microlensing (this assumption can also be adopted, to some extent, for low ionization broad emission lines; \\citealp{2000ApJ...533..631K}; \\citealp{2002ApJ...576..640A}). { If we define the baseline from emision lines measured in the same wavelength regions as the continua affected by microlensing, we can also remove the extinction and isolate the microlensing effects.  ``Intrinsic'' flux ratios between the images in the absence of microlensing can be determined from the observation of the mid-infrared and radio-emitting regions of quasars that should also be large enough to average out the effects of microlensing (see \\citealp{2004ApJ...605...58K} and references therein). However, the extinction at mid-infrared and radio wavelengths is lower than the extinction at the wavelengths in which microlensing is usually detected and measured (optical, near-infrared, and X-ray). Consequently, the difference between the mid-infrared (radio) and the optical (X-ray or near-infrared) continuum fluxes will include not only the effects of microlensing but also the effects of extinction. In addition, note that the availability of data at optical wavelengths is considerably greater than at other wavelengths.}  Thus { we will use the NEL and} continuum flux ratios among the different images of a lensed QSO to estimate the difference of microlensing magnification between the images at a given epoch with certain restrictions that we detail in the following paragraphs.  The flux (in magnitudes) of an emission line observed at wavelength $\\lambda$ of image $i$ of a multiply imaged quasar is equal to the flux of the source, $m_0^{lin}\\left({\\lambda\\over 1 +z_s}\\right)$, magnified by the lens galaxy (with a $\\Phi_i$ magnification factor; $\\mu_i=-2.5\\log \\Phi_i$) and corrected by the extinction of this image caused by the lens galaxy, $A_i\\left({\\lambda\\over 1 +z_l}\\right)$ (see, e.g., \\citealp{2004ApJ...605..614M}),  \\begin{equation} m_i^{lin}(\\lambda)= m_0^{lin}\\left({\\lambda\\over 1 +z_s}\\right) + \\mu_i + A_i\\left({\\lambda\\over 1 +z_l}\\right), \\end{equation}  \\noindent where $z_s$ and $z_l$ are the redshifts of the source and the lens, respectively.  In the case of the continuum emission, we must also take into account the intrinsic variability of the source combined with the delay in the arrival of the signal, which is different for each image, $\\Delta t_i$, and the microlensing magnification, which depends on wavelength and time (with a $\\phi_i\\left[{\\lambda\\over 1 +z_l},t \\right]$ magnification factor; $\\Delta\\mu_i=-2.5\\log \\phi_i$),  \\begin{equation} m_i^{con}(\\lambda,t)= m_0^{con}\\left({\\lambda\\over 1 +z_s},t-\\Delta t_i\\right) + \\mu_i + A_i\\left({\\lambda\\over 1 +z_l}\\right) + \\Delta\\mu_i \\left({\\lambda\\over 1 +z_l},t\\right). \\end{equation}  Thus, the difference between continuum and line fluxes { cancels the terms corresponding to intrinsic magnification and extinction ($\\mu_i + A_i$)}:  \\begin{equation} m_i^{con}(\\lambda,t) - m_i^{lin}(\\lambda)= m_0^{con}\\left({\\lambda\\over 1 +z_s},t-\\Delta t_i\\right) - m_0^{lin}\\left({\\lambda\\over 1 +z_s}\\right) + \\Delta\\mu_i\\left({\\lambda\\over 1 +z_l},t\\right). \\end{equation}  If we consider a pair of images, $1$ and $2$, the continuum ratio relative to the zero point defined by the emission line ratio  can be written (in magnitudes) as,   \\begin{eqnarray} \\Delta m (\\lambda,t)= (m_1-m_2)^{con}-(m_1-m_2)^{lin}=&{\\Delta\\mu_1\\left({\\lambda\\over 1 +z_l},t\\right)- \\Delta\\mu_2\\left({\\lambda\\over 1 +z_l},t\\right) } \\label{eq_baseline} \\\\ &+\\Delta m_0^{con}\\left({\\lambda\\over 1 +z_s},\\Delta t_1-\\Delta t_2\\right).\\nonumber \\end{eqnarray}   We have referred the equations for the magnification of both images to an arbitrary time, $t$ (note that microlensing-induced variability between a pair of images is uncorrelated).  The first term of Equation (\\ref{eq_baseline}) is the relative microlensing magnification between images 1 and 2.  The significance of the second term, $\\Delta m_0^{con}$, which represents the source variability, can be estimated by comparing the intrinsic quasar variability on time-scales typical of the time delay between images in gravitational lens systems with the expected distribution of microlensing magnifications. As we shall  discuss in Section \\ref{sect_conclusions}, the intrinsic source variability is not significant for our computations.  In summary, with the proposed method similar information as in the Milky Way MACHO experiments is obtained but with a single-epoch measurement. The objective of this study is to apply this method to published data of quasar microlensing. In \\S \\ref{sect_observ} we collect the data from the literature and fit models to the systems of multiply imaged quasars to derive suitable values of the projected matter density at the image locations. Using these values, probabilistic models for microlensing magnifications are derived in Section \\ref{sect_maps} for a range of fractions of mass in compact objects. Sections \\ref{sect_mle} and \\ref{sect_scaling} are devoted to estimate this fraction. Finally, in \\S \\ref{sect_conclusions} we present and discuss the main conclusions. ",
        "conclusions": "}  In the previous sections we have extended to the extragalactic domain the local (Milky Way, LMC, and M31) use of microlensing to probe the properties of the galaxy halos. Although our primary aim was to explore the practical application of the proposed method, we found that the data available in the literature can be consistently interpreted only under  the hypothesis of a very low mass fraction in microlenses; { at 90\\% confidence: $\\alpha(\\log L_{max})=0.05_{-0.03}^{+0.09}$ (maximum likelihood estimate)}  {for a quasar continuum source of intrinsic size  $r_s=2.6 \\cdot 10^{15}\\rm\\, cm$}. This result arises directly from the shape of the histogram of microlensing magnifications, with a maximum of events close to no magnification and stands for a wide variety of microlensing models statistically representative of the considered image pairs. There is a dependence of the estimate of $\\alpha$ on the source size but high values of the mass fraction ($\\alpha>0.2$) are possible only for unexpectedly large source sizes ($r_s>4 \\cdot 10^{16}\\rm\\, cm$). The low mass fraction is in good agreement with the results of EROS \\citep{{2007A&A...469..387T}} for the Milky Way and with the limit established by MEGA \\citep{2006A&A...446..855D} for M31. { The agreement is also good with the few microlensing-based estimates available for individual objects. In RXJ 1131-1231 \\citet{2009arXiv0906.4342D} found $\\alpha\\sim 0.1$. In PG 1115+080 \\citet{2008ApJ...689..755M} obtained values in the range $\\alpha=0.08$ to $0.15$. For the same system \\citet{2009ApJ...697.1892P} found $\\alpha\\sim 0.1$ for a source of size $r_s=1.3 \\cdot 10^{16}\\rm\\, cm$}.  {On the other hand, our estimate of the fraction of mass in microlenses, $\\alpha(\\log L_{max})=0.05_{-0.03}^{+0.09}$, approximates the expectations for the fraction of visible matter. {  \\cite{{2007ApJ...671.1568J}}, for instance, comparing the mass inside the Einstein ring in 22 gravitational lens galaxies with the mass needed to produce the observed velocity dispersion, inferred average stellar mass fractions of $0.026\\pm 0.006$ (neglecting adiabatic compression) and $0.056\\pm 0.011$ (including adiabatic compression). As discussed in \\cite{{2007ApJ...671.1568J}} these values are also in agreement with other estimates of the stellar mass fraction that relied on stellar population models: $\\sim 0.08$ \\citep{2006ApJ...648..826L},  $0.065_{-0.008}^{+0.010}$ \\citep{2005ApJ...635...73H}, and $0.03_{-0.01}^{+0.02}$ \\citep{2006MNRAS.368..715M}. Thus, we can conclude} that microlensing is { probably} caused by stars in the lens galaxy, and that there is no statistical evidence for MACHOS in the halos of the 20 galaxies of the sample we considered.  How robust are these results? There are several sources of uncertainty to consider.  Firstly, we neglected in Equation \\ref{eq_baseline} the term arising from source variability, $\\Delta m_0^{con}$. From  a group of 17 gravitational lenses with photometric monitoring available in the literature we estimate an average gradient of variability of $0.1\\rm\\, mag\\, year^{-1}$. Taking into account that the average delay between images is about $3\\rm\\, months$ (a conservative estimate; note that the group of  lens systems used includes many doubles, some of them with very large time delays) we can expect an amplitude related to intrinsic source variability of $\\Delta m_0^{con}\\sim 0.03$, which, according to the histogram of magnifications (Figure \\ref{fig_histo_A}), is not significant. Moreover, if we assume that the probability of $\\Delta m_0^{con}$ is normally distributed, the global effect of source variability is to broaden the histogram of microlensing magnifications, diminishing the peak and enhancing the wing. In other words, source variability leads to an overestimate of $\\alpha$. Thus, the mass fraction should be even lower if significant source variability were hidden in the data. In the same way, other sources of error in the data, such as the difficulty in separating line and continuum or in removing from the narrow emission lines the high ionization broad emission lines that could be partially affected by microlensing, probably tend to induce additional magnitude differences, $\\Delta m$, between the images and, hence, to an overestimate of $\\alpha$. On the contrary, cross-contamination between the spectra of a pair of images masks the impact of microlensing and may affect our results. Although most of the bibliographic sources of microlensing measurements analyze this problem concluding that the spatial resolution was sufficient to extract the spectra without contamination, it is clear that high S/N data obtained in subarcsecond seeing conditions will help to control this important issue.  Another point to address is the treatment of some of the quads, where only a subset of the images are used. Are we systematically excluding faint images that might be highly demagnified by microlensing? Let us examine the four incomplete quads in our sample. The fold lens SDSS J1004+4112 has two close images A and B. A is probably a saddle-point image and shows the most anomalous flux (Ota et al. 2006). In contrast, the optical/X-ray flux ratios of C and D are almost the same. Thus, there is no reason to suppose that the image without a useful spectrum (D) has higher microlensing probability than the others. PG 1115+080 is another fold quad. $A_1$ and $A_2$ are the two images closest to the critical curve and have a (moderately) anomalous flux ratio and optical variability (Pooley et al. 2007). The two images without available spectra (C and D) show only a  small optical variability and are not particularly prone to microlensing. In RXS J1131-1231 the most anomalous flux ratio is B/C and A is a saddle-point image (Sluse et al. 2006). Image D (the one with no available spectrum)  also has an anomalous flux but is not more susceptible to microlensing than the other images. Thus, in three of the four incomplete quads there is no reason to suppose that we are biasing the sample towards image pairs with lower microlensing probability. The case of SDSS 0924+0219 is more problematic. There are two sets of data for this  object, one by Eigenbrod et al. (2006) based on observations of the low ionization lines MgII and CIII], which, after two epochs of observation, reveals no difference between the line and continuum flux ratios of components A and B. The other set of data (Keeton et al. 2006) is based on Ly$\\alpha$ observations (a high ionization emission line suposed to come from a smaller region than the low ionization emission lines) and microlensing is detected not only in the continuum but also in the emission lines. This implies that the baseline for no microlensing magnification cannot be defined and, consequently, we could not\u00a0 consider Keeton et al. (2006) results. Anyway, we have repeated (as a test)  the entire maximum likelihood estimate procedure \u00a0to  derive $\\alpha$ but now using for SDSS 0924+0219 the microlensing measurements by Keeton et al. (2006). The results are almost identical: $\\alpha = 0.05^{+0.10}_{-0.03}$.  The size of the sample also limits the statistical interpretation. An improvement in the S/N of the histogram of microlensing magnifications is very important to ascertain the statistical significance of the low frequency of events at large magnification (only two events of high magnification are detected), which can impose severe constraints on the microlensing models. Another reason to increase the size of the sample is the possibility to define subsamples at different galactocentric distances where different ratios of visible to dark matter are expected. In the same way it would be possible to define subsamples according to the type of lens galaxy or other interesting properties of lens systems.  In any case, the impact of the main result of our study $-$absence of MACHOS in the 10 to 0.1$M_\\odot$ mass range in the halos of lens galaxies$-$ and its future prospects, points to the need to improve the statistical analysis in two ways: increasing the number, quality and homogeneity of the microlensing magnification measurements from new observations, and reducing the uncertainties in the macro-lens models."
    },
    "0910/0910.0844_arXiv.txt": {
        "abstract": "Most of the super massive black hole mass (\\Mbh) estimates based on stellar kinematics use the assumption that galaxies are axisymmetric oblate spheroids or spherical. Here we use fully general triaxial orbit-based models to explore the effect of relaxing the axisymmetric assumption on the previously studied galaxies \\M32 and \\NGCN{3379}. We find that \\M32 can only be modeled accurately using an axisymmetric shape viewed nearly edge-on and our black hole mass estimate is identical to previous studies. When the observed 5\\dgr\\ kinematical twist is included in our model of \\NGCN{3379}, the best shape is mildly triaxial and we find that our best-fitting black hole mass estimate doubles with respect to the axisymmetric model. This particular black hole mass estimate is still within the errors of that of the axisymmetric model and consistent with the \\Msigma\\ relationship. However, this effect may have a pronounced impact on black hole demography, since roughly a third of the most massive galaxies are strongly triaxial. ",
        "introduction": "The masses of super massive black holes in the centers of galaxies are known to correlate with several properties of the host galaxy. The most well known correlation is with the stellar velocity dispersion of the galaxy \\citep[\\Msigma, e.g.,][]{2002ApJ...574..740T}. The black hole is thought to play an important role in the evolution of its host (e.g. AGN feedback), as the properties of galaxies are tightly linked to \\Mbh. Therefore it is important to be able to measure the \\Mbh\\ accurately and understand whether the scatter in the relations is due to measurement error, or if it is intrinsic. Most of the \\Mbh\\ estimates upon which these relationships are based were derived using dynamical (edge-on) axisymmetric (or spherical) dynamical models (See \\citealp{2005SSRv..116..523F} for a review).  It has long been known from photometry that some elliptical galaxies are triaxial \\citep[e.g.][]{1996ApJ...464L.119K}, i.e.\\ have intrinsic shapes with three distinct axes \\citep{1976MNRAS.177...19B, 1978MNRAS.183..501B}, and more recently \\cite{2007MNRAS.379..401E} and \\cite{2007MNRAS.379..418C} have shown using stellar kinematics that around a third of the most massive ellipticals are at least mildly triaxial. It is thus relevant to rederive the \\Mbh\\ in these galaxies with our triaxial instead of axisymmetric orbit super-position models.  \\cite{2007MNRAS.381.1672T} have modelled mock triaxial merger remnants using axisymmetric geometry and found a correlation between viewing angle and the recovered total galaxy mass. Additionally, they found that the luminous mass to light ratio (\\ML) was underestimated by up to 20\\% confirming that triaxial modelling is preferred.  In this paper we explore the black hole recovery with the triaxial machinery from \\cite{2008MNRAS.385..647V}\\ using galaxies that have previously been modeled with axisymmetric codes to verify the triaxial models. We do this with \\M32\\ and \\NGCN{3379} which have their black hole mass determined using \\Stis, as well as \\Sauron~and \\Oasis~Integral Field Unit (IFU) data. Both galaxies have state-of-the-art kinematics available over a large spatial range and are inside the sphere of influence of the black hole.  We describe our modeling technique and uncertainties in \\S\\ref{method} and then derive our black hole estimates on galaxies \\M32\\ and \\NGCN{3379}\\ in \\S\\ref{sec:bh3}. We briefly address the reliability in \\S\\ref{sec:tbh_reliability}, and we end with discussion and conclusions in \\S\\ref{sec:tbh_discuss}. ",
        "conclusions": "\\label{sec:tbh_discuss} %  We have shown two applications of the triaxial orbit super-position method to galaxies that had their central black hole mass measured with axisymmetric models. The reason for this was two-fold: First, we confirmed that we obtain the same \\Mbh\\ when using our new triaxial implementation in the oblate axisymmetric limit, confirming that our method is reliable. Secondly, we explored what might happen when the assumption of axisymmetry is relaxed.  For the nearby elliptical galaxy \\M32\\ we obtained identical results to previous axisymmetric modeling finding a \\Mbh\\ of ($2.4\\pm1.0) \\times10^{6}$ \\Msun. Our best-fitting shape of \\M32\\ is very close to oblate axisymmetric, excluding strongly triaxial shapes. The viewing angles are not well constrained.  We also duplicated the models of \\NGCN{3379}. After confirming the results from \\citetalias{2006MNRAS.370..559S} in the axisymmetric limit, we expanded our search to include triaxiality. While the shape is not strongly constrained, we found that this galaxy is seen almost face-on and can best be described with a triaxial model. The best-fitting models are very round in the center and become fairly flattened at larger radii. These results confirm the claims about the intrinsic shape from \\cite{1991ApJ...371..535C}, \\cite{1994AJ....108..111S} and \\cite{2008MNRAS.390...93K}.   The black hole in this triaxial model weighs $(4\\pm 1) \\times 10^{8}$\\Msun, which is two times bigger than the axisymmetric estimate from \\citetalias{2006MNRAS.370..559S} and \\cite{2000AJ....119.1157G}. We speculate that the difference in the estimate is due to the combined effect of changing from an edge-on to a nearly face-on view and the inclusion of the strongly radial box orbits in the modeling.%  The significance of the change of the \\Mbh\\ of \\NGCN{3379} is unclear. While this individual result is still within the scatter of the \\Msigma\\ relation, it can greatly affect all empirical relations based on \\Mbh, like \\Msigma, if similar effects are seen in more galaxies. If many intrinsically triaxial galaxies would have their black hole mass underestimated using axisymmetric modeling, the empirical relations would be systematically underestimating black hole masses at the top end, as we believe that triaxial galaxies dominate at the high mass end \\citep[e.g.][]{1996ApJ...464L.119K, 2007MNRAS.379..401E}. To study this in more detail, it is necessary to study the nuclei of the most massive galaxies with triaxial models.  Also, the axisymmetric models and their \\Mbh\\ estimates have only been tested with purely axisymmetric (and spherical) test models. This means that if axisymmetric models cannot accurately recover black hole masses in triaxial galaxies, and if most galaxies are significantly triaxial, then this would also have an impact on black hole demography\\footnote{Of course, a similar argument could be held for the triaxial models, if the modeled galaxies cannot be described by stable triaxial geometries.}. A strong hint in this direction is given by \\cite{2007MNRAS.381.1672T}, who showed that the axisymmetric models have difficulty estimating the \\ML\\ in triaxial galaxies. Based on their results and our \\NGCN{3379} model, we speculate that the \\Mbh\\ estimates in triaxial galaxies derived using axisymmetric modeling will have systematic errors, because the recovery of the \\Mbh\\ is strongly linked to the \\ML."
    },
    "0910/0910.4240_arXiv.txt": {
        "abstract": "While there have been multiple observational programs aimed at detecting linear polarization of optical radiation emitted by  ultracool dwarfs, there has been comparatively less rigorous theoretical analysis of the problem. The general expectation has been that the atmospheres of those substellar-mass objects with condensate clouds would give rise to linear polarization due to scattering. Because of rotation-induced non-sphericity, there is expected to be incomplete cancellation of disk-integrated net polarization and thus a finite polarization. For cloudless objects, however, only molecular Rayleigh scattering will contribute to any net polarization and this limit has not been well studied.  Hence in this paper we present a detailed multiple scattering analysis of the polarization expected from those T-dwarfs whose spectra  show absence of condensates. For this, we develop and solve the full radiative transfer equations for linearly polarized radiation. Only atomic and molecular Rayleigh scattering  are considered to be the source of polarization. We compute the local polarization at different angular directions in a plane-parallel atmospheres calculated for the range of effective temperatures of T dwarfs and then average over the whole surface of the object. The effects of gravity and limb darkening as well as rotation  induced non-sphericity are included. It is found that the amount of polarization decreases with the increase in effective temperature. It is also found that significant polarization at any local point in the atmosphere arises only in the optical (B-band). However, the disk integrated polarization--even in the B-band--is negligible.  Hence we conclude that, unlike the case  for cloudy L dwarfs, polarization of cloudless T-dwarfs by atomic and molecular  scattering may not be detectable.  In the future we will extend this work to cloudy L- and T-dwarf atmospheres. ",
        "introduction": "Brown dwarfs have masses sufficient to ignite deuterium burning but insufficient to enter into the hydrogen burning main sequence. Instead, they radiate away the gravitational potential energy produced during their formation and contraction.  As an individual brown dwarf cools it progressively passes through spectral types ranging from late M through the L and T sequences.  Helpful reviews of brown dwarf science may be found in  Chabrier \\& Baraffe (2000), Burrows et al. (2001), Kirkpatrick (2005), and Marley \\& Leggett (2008).  While condensates such as iron, forsterite ($\\rm Mg_2SiO_4$), and enstatite ($\\rm MgSiO_3$) are favored to be formed in the atmosphere of relatively warm L dwarfs \\citep{lunine86, tsuji96, lodders99, burrows00, allard01}, in the atmosphere of mid-type and later T dwarfs, grains condense well bellow the photosphere and are not an important opacity source \\citep{ackerman01, tsuji02, chabrier00, burrows01, kirkpatrick05, marley08}. Synthetic spectra of model atmospheres lacking condensates fit well with the observed spectra  of field T dwarfs later than about type T4 (with effective temperature $T_{\\rm eff}\\sim 1200\\,\\rm K$; \\cite{stephens09}) suggesting the absence of condensates in the visible atmosphere of these objects. At the same time, inclusion of condensate cloud of various species explains the spectra of L dwarfs (e.g., Cushing et al. 2008).  \\cite{sengupta01} first predicted that scattering of light by the condensates in the photopshere of L dwarfs would give rise to  a significant amount of linear polarization. This polarization would of course cancel out when integrated over a spherical surface of the star. However, analysis of high resolution spectra of L dwarfs confirms that these objects are rapid rotator with the projected rotational velocity ranging from as high as 90 to 10 $\\rm kms^{-1}$ \\citep{basri00,mohanty03, osorio06, reiner, jones}.  A rapid rotation around its axis induces departure from sphericity in the shape of the object.  As a consequence, the net polarization integrated over the disk does not completely cancel out. Subsequently, linear polarization of several L dwarfs has been detected by various group of astronomers \\citep{menard02, osorio05, goldman09, tata09}. The observation were analyzed and explained by \\cite{sengupta03, sengupta05} who assumed a single scattering atmosphere wherein the effect of atoms and molecules was neglected. These authors suggested that scattering by  atoms and molecules would give rise to negligible amount of polarization and  hence polarization of T dwarfs in the optical should be inconsequential. However their single scattering approach left room for some uncertainty because they  considered a purely scattering medium while a proper approach would incorporate both absorption and multiple scattering.  In this paper we present the results of our detailed investigation on the polarization of T-dwarfs caused by multiple scattering of atoms and molecules in cloud-less atmospheres. In the next section we discuss the atmospheric models. In section~3, we present in detail the equations of transfer of polarized radiation and in section~4 we present the numerical methods of solving the transfer equation of polarized radiation in a plane-parallel medium. The local polarization is integrated over the rotationally  distorted surface.  We briefly describe the method adopted to calculate the rotational oblateness and the methods of integrating the polarization over a rotationally distorted stellar disk in section~5. We discuss the results in section~6 followed by our conclusion in the last section. ",
        "conclusions": "In the present paper, we have discussed the results of our investigation on multiple scattering polarization of cloudless T dwarfs. We predict that in the absence of condensates in the photosphere of T dwarfs later than about type T3 \\citep{stephens09}, non-zero polarization would arise only at wavelengths shorter than ${\\rm 0.6 \\mu m}$- a spectral region where these dwarfs are exceptionally dark. In order to calculate the maximum possible polarization of T dwarfs, we have considered a polytropic equation of state with index $n=1$ which is the minimum  permissible value of $n$ for representing the matter distribution in brown dwarfs. Consequently, the rotation-induced oblateness is maximized for a given rotational velocity.  We find that even with a rotational velocity as high as 90 kms$^{-1}$, the maximum amount of disk integrated polarization that arises at an inclination angle of $90^o$, is of the order $10^{-4}$, too small to be detected  in the B-band.  The degree of polarization would be greatest for the coolest objects and it would decrease with increasing  effective temperature.  The results support the claim by \\cite{sengupta03, sengupta05} that the polarization detected in the optical region (R- and I-band) of L dwarfs has negligible contribution from atomic and molecular scattering. Therefore, if linear polarization in R- or I-band of T dwarfs is detected in future, it will imply the presence of condensates in the photosphere of T dwarfs.  Although a magnetic field in principle might produce polarization,  T dwarf atmospheres are highly neutral (Gelino et al. 2002) and thus the presence of magnetic field  would not give rise to detectable amount of polarization in the optical \\citep{menard02}. Finally, polarization could play a crucial role in measuring the loss of condensates from brown dwarf atmospheres as we expect L- and early T-type dwarfs to show linear polarization in R- and I-band by dust scattering in the phototosphere whereas mid-type and later T dwarfs would not show any polarization at wavelength greater than ${\\rm 0.6 \\mu m}$. The detailed investigation of the polarization of L  dwarfs by multiple dust scattering will be presented in a forthcoming paper."
    },
    "0910/0910.2699_arXiv.txt": {
        "abstract": "{Recent attempts to constrain cosmological variation in the fine structure constant, $\\alpha$, using quasar absorption lines have yielded two statistical samples which initially appear to be inconsistent. One of these samples was subsequently demonstrated to not pass consistency tests; it appears that the optimisation algorithm used to fit the model to the spectra failed. Nevertheless, the results of the other hinge on the robustness of the spectral fitting program VPFIT, which has been tested through simulation but not through direct exploration of the likelihood function. We present the application of Markov Chain Monte Carlo (MCMC) methods to this problem, and demonstrate that VPFIT produces similar values and uncertainties for $\\Delta\\alpha/\\alpha$, the fractional change in the fine structure constant, as our MCMC algorithm, and thus that VPFIT is reliable. ",
        "introduction": "Recent years have seen sustained interested in attempting to determine whether any of the fundamental constants of nature vary. Efforts have been focused in particular on the fine structure constant, $\\alpha$, and the proton-to-electron mass ratio, $\\mu$, although other dimensionless constants have also been considered. Quasars provide a method of probing the value of $\\alpha$ in the early universe by illuminating gas clouds along the line of sight to Earth. In particular, certain absorption lines in these gas clouds are sensitive to changes in one or more of the fundamental constants of nature. For an atom/ion, the relativistic corrections to the energy levels of an electron are proportional to $\\alpha^2$, although the magnitude of the change depends on the transition under consideration. By comparing the wavelengths of absorption lines with differing sensitivities to a change in $\\alpha$, we are able to place constraints on a change in $\\alpha$ between the early universe and today. That is, we are able to measure $\\Delta\\alpha/\\alpha = (\\alpha_z-\\alpha_0)/\\alpha_0$, where $\\alpha_z$ is the value of $\\alpha$ at redshift $z$ and $\\alpha_0$ is the laboratory value.  The tentative detection of a variation in $\\alpha$ reported by \\citet{Webb99} has further increased interest in this field. Subsequent efforts refined this work by increasing the number of absorption systems considered. This yielded a constraint of $\\Delta\\alpha/\\alpha = (-0.57 \\pm 0.11)\\times 10^{-5}$ \\citep{Murphy04} from 143 absorption systems. However all of these observations are from the Keck/HIRES (High Resolution Echelle Spectrometer) instrument, and so it remains important to confirm such unexpected results with independent equipment.  \\citet{Chand04} reported $\\Delta\\alpha/\\alpha = (-0.06 \\pm 0.06) \\times 10^{-5}$ from 23 measurements using the Ultraviolet and Visual Echelle Spectrograph (UVES), on the Very Large Telescope (VLT). These two measurements of $\\Delta\\alpha/\\alpha$ are clearly discrepant. A later analysis \\citep{Murphy08a} found that the analysis of \\citeauthor{Chand04} cannot be correct as it states a statistical precision in excess of the maximum theoretical precision allowed by the data. Similarly, \\citeauthor{Murphy08a} analysed the $\\chi^2$ vs $\\Delta\\alpha/\\alpha$ curves produced by the \\citeauthor{Chand04} optimisation algorithm, and concluded that the shape of the curves demonstrate a failure of the algorithm to find the maximum likelihood value of $\\Delta\\alpha/\\alpha$ implied by the model fits and data, and thus that the estimate of $\\Delta\\alpha/\\alpha$ given by \\citeauthor{Chand04} is unreliable.  Although it would appear that the \\citeauthor{Murphy04} results are robust, it is worth directly investigating the optimisation algorithm used in order to confirm that it is reliable. Furthermore, each measurement of $\\Delta\\alpha/\\alpha$ will be subject to systematic errors, but some systematic errors should have an expectation value of zero, and thus averaging over many absorption systems will in principle eliminate such errors. It has become commonplace to quote values of $\\Delta\\alpha/\\alpha$ for individual systems; in these cases, one has no method of determining the size of certain systematic errors through comparison with other systems, and so one should be particularly careful that the statistical errors are stated correctly. The MCMC method we describe herein allows one to confirm the validity of the statistical errors produced by a different method of analysis. ",
        "conclusions": "Given suitable knowledge of the observed distribution of column densities and Doppler parameters, we could implement these distributions as priors to our model. However, our statistical constraints are generally sufficiently good that this should not significantly alter our parameter estimates. In any event, we are primarily interested in $\\Delta\\alpha/\\alpha$, for which a flat prior is reasonable on ignorance grounds.  We would like to apply our algorithm to verify the results of \\citet{King08} in relation to $\\Delta\\mu/\\mu$; however, those fits involve greater than a thousand parameters each. In this context, our algorithm is hopelessly inadequate. Our running times for the objects described herein varied from a few hours to a few days, and have only a few tens of parameters. Exploration of more complicated cases must wait for advances in computing power.  Our results demonstrate that VPFIT produces reliable parameters estimates and uncertainties for relatively simple situations. Experience with VPFIT suggests that there does not appear to be any indication of failure with moderately complicated circumstances, and so we would argue that the optimisation algorithm used by VPFIT is robust. The implication of is this it that the results of \\citet{Murphy04} are unlikely to be explained by some failure of the optimisation algorithm used by VPFIT. Thus the detection of a change in $\\alpha$ must either be real, or due to some other unknown issue."
    },
    "0910/0910.5246_arXiv.txt": {
        "abstract": "We examine large-scale fluctuations in the \\ion{He}{2} \\lya forest transmission during and after \\ion{He}{2} reionization.  We use a simple Monte Carlo model to distribute quasars throughout a large volume and compute the resulting radiation field along one-dimensional skewers.  In agreement with previous studies, we find that the rarity of these sources induces order unity fluctuations in the mean optical depth after reionization, even when averaged over large segments ($\\sim 10$--$100 \\Mpc$ across).  We compare our models to existing data along five \\ion{He}{2} \\lya forest lines of sight spanning $z \\sim 2$--$3.2$.  The large cosmic variance contained in our model plausibly explains many of the observed fluctuations at $z \\la 2.5$.  But our models cannot accommodate the large fluctuations toward high optical depths on $\\sim 30 \\Mpc$ scales observed at $z \\sim 2.7$--$2.9$, and the measured optical depths ($\\tau_{\\rm eff} \\ga 4$) at $z>2.9$ are difficult to explain with a smoothly-evolving mean radiation field.  In order to better understand this data, we construct a toy model of \\ion{He}{2} reionization, in which we assume that regions with the smallest radiation fields in a post-reionization Universe (or farthest from strong ionizing sources) are completely dark during reionization.  The observed fluctuations fit much more comfortably into this model, and we therefore argue that, according to present data, \\ion{He}{2} reionization does not complete until $z \\la 2.9$. ",
        "introduction": "\\label{intro}  Our premier tool for studying the intergalactic medium (IGM) is the \\lya forest, from which we can measure the absorption properties of the low-density ``cosmic web\" lying between galaxies.  Although observations of the \\ion{H}{1} forest have a long and rich forty-year history \\citep{rauch98}, recently there has been great progress in measuring absorption from other ions, including both heavy elements and helium, the second most common element in the Universe.  Unfortunately, observing the \\ion{He}{2} \\lya forest is difficult, because the transition's rest wavelength of 304 \\AA \\ places much of the absorption in the far ultraviolet, even for material at $z \\sim 2$--$3$.  Not only do ultraviolet telescopes present formidable technical challenges, but studying the forest also requires a luminous far-UV source to provide a background light.  Such lines of sight are rare, because any that intersect strong \\ion{H}{1} absorbers do not retain their UV flux.  To date, five quasars at $z \\sim 2$--$3.5$ have satisfied these criteria and have high-resolution \\ion{He}{2} \\lya forest measurements (see \\S \\ref{data} below for references to the individual quasars), although very recently more than twenty other potential targets have been identified \\citep{syphers09, syphers09b}.  Nevertheless, the known lines of sight are particularly interesting, because current data (from these measurements and other indirect techniques) suggest that quasars fully ionized \\ion{He}{2} by $z \\sim 3$.  Most obviously, there is a substantial decrease in the mean \\lya forest absorption beyond $z \\sim 2.9$ \\citep{dixon09}, which is compatible with models where large swathes of \\ion{He}{2} are cleared out by quasar light at about that time \\citep{sokasian02, furl08-helium, mcquinn09}.  Indirect evidence from the IGM temperature also suggests reionization at $z \\sim 3.2$--$3.4$ \\citep{schaye00, ricotti00, lidz09-temp}, which is consistent with models of helium reionization ending at these slightly higher redshifts \\citep{gleser05, furl08-igmtemp, mcquinn09}.  Here, we will study the variability in the \\ion{He}{2} \\lya forest data and its consequences for reionization.  Specifically, recent work has shown that the high-energy radiation background fluctuates strongly both during and after helium reionization, thanks largely to the rarity of the bright quasars providing the ionizing photons (see \\citealt{fardal98, maselli05, bolton06, meiksin07, tittley07, furl09-hefluc} for theoretical work and \\citealt{shull04, fechner06, fechner07} for observations).  Moreover, these fluctuations can span large physical scales \\citep{furl09-heps}.  We might therefore expect that cosmic variance from the small number of well-studied lines of sight may limit any firm conclusion based on the optical depth.  On the other hand, the existing observations already exhibit substantial fluctuations and may themselves suggest something useful about reionization.  We naturally expect larger fluctuations during reionization, because the radiation field can be extremely large near quasars but almost zero in the \\ion{He}{2} regions far from these sources \\citep{meiksin07, furl09-hefluc}.  Here, we will quantify how these fluctuations impact measurements of the mean \\lya forest absorption and show that the observed variations are by themselves powerful probes of the radiation field.  We will use a simple Monte Carlo model and ask whether the existing data are consistent with \\ion{He}{2} reionization ending at $z \\sim 3$.  Our approach is similar to \\citet{fardal98}, although we use updated input parameters, add a toy model for reionization itself, and have the advantage of comparing to significantly more data.  This paper is organized as follows.  In \\S \\ref{mc}, we describe the Monte Carlo model that we use for the post-reionization limit.  We present a simple toy model of \\ion{He}{2} reionization in \\S \\ref{reion} and illustrate the basic characteristics of these models in \\S \\ref{pdfs}.  In \\S \\ref{data}, we compare our results to existing data.  Finally, we conclude in \\S \\ref{disc}.  In our numerical calculations, we assume a cosmology with $\\Omega_m=0.26$, $\\Omega_\\Lambda=0.74$, $\\Omega_b=0.044$, $H=100 h \\hunits$ (with $h=0.74$), $n=0.95$, and $\\sigma_8=0.8$, consistent with the most recent measurements \\citep{dunkley09}.  Unless otherwise specified, we use comoving units for all distances. ",
        "conclusions": "\\label{disc}  We have computed the fluctuations expected in measurements of the average transmission of the \\ion{He}{2} \\lya forest due to large-scale variations in the radiation field. We used a simple Monte Carlo model of the quasar distribution to simulate the radiation field.  We then measured $\\tau_{\\rm eff}$ averaged over long ($\\sim 30$--100 Mpc) segments along random lines of sight and computed its probability distribution.  For a fully-ionized Universe, the radiation field typically produces a full-width at half-maximum for the $\\tau_{\\rm eff}$ distribution of order its mean value \\citep{fardal98}.  We also constructed a toy model of reionization, where we demanded that a fraction $x_{\\rm dark}$ of pixels have zero transmission.  We chose these dark pixels by identifying the regions that would have the smallest transmission in a post-reionization Universe.  Such a model is reasonable, because only active quasars can illuminate a region strongly enough to render it transparent, even if fossil, mostly-ionized regions fill much of the IGM \\citep{furl08-fossil}. This method produces long dark stretches (with sizes $\\ga 100 \\Mpc$) in the empty voids between quasars, even late in reionization.  This procedure broadens the $\\tau_{\\rm eff}$ distribution by producing a large tail toward very high optical depths (and, to compensate, shifting the median transmission to smaller values).  We then compared our model to the observed fluctuations in the existing sample of \\ion{He}{2} \\lya forest lines of sight, as compiled by \\citet{dixon09}.  In particular, we asked whether a post-reionization model with a smoothly evolving attenuation length (and hence no change in the ionization state of IGM helium) could account for the observed data.  At lower redshifts ($z \\sim 2.25$) such a model is reasonable, even though the data show somewhat stronger absorption than expected.  The observed transmission is unlikely according to our fiducial model, but modest changes to the normalization or slope of the underlying $\\VEV{\\Gamma}$ evolution can easily accommodate it.  Alternatively, the underlying IGM density fluctuations may also account for the discrepancy \\citep{fardal98}.  This points to one important (though unsurprising) lesson from our analysis: even with the modest signal-to-noise already available in \\ion{He}{2} \\lya forest spectra, the statistical errors on $\\tau_{\\rm eff}$ are significantly smaller than the natural level of cosmic variance.  Thus, it is crucial to consider the latter when interpreting the existing data.  At higher redshifts, the data become much more difficult to explain.  Even though there are only five lines of sight with existing measurements, we found strong evidence that a smoothly-evolving post-reionization model is \\emph{inconsistent} with the observations.  At $z \\sim 2.7$--$2.9$, the measurements fluctuate strongly toward high optical depths on $\\sim 20$--$35 \\Mpc$ scales.  Such regions of deep absorption cannot fit into a post-reionization model (at $\\ga 99\\%$ confidence), because the large number of quasars far from the line of sight already provide enough radiation to keep the ionized fraction reasonably large.  Thus, the deviation cannot simply be explained as cosmic variance from a single line of sight, unless it is truly pathological.  To reconcile the observations and model, we need to increase the predicted $\\VEV{\\tau_{\\rm eff}}$ significantly, use relatively small mean free paths, and take a conservative view of the data (i.e., accepting the more conservative \\citealt{kriss01} lower limits).  However, in this case there \\emph{still} must be a break in $\\VEV{\\tau_{\\rm eff}}(z)$ at $z \\sim 2.7$, which we have previously argued would likely itself signal \\ion{He}{2} reionization \\citep{dixon09}.  At $z \\ga 2.9$, the only well-observed line of sight (Q0302-003) has $\\tau_{\\rm eff} \\ga 4$ in two separate $\\sim 60$--$100 \\Mpc$ intervals.  This opacity is much larger than our smoothly-evolving model predicts, and even when post-reionization radiation fluctuations and density fluctuations are included, it remains extremely improbable.  Thus, this line of sight on its own also argues strongly for either ongoing reionization or a rapid increase in $\\VEV{\\tau_{\\rm eff}}$ at $z \\sim 2.9$, likely also signaling the end of reionization.  Our model, therefore, supports the conclusion of \\citet{dixon09}, who examined the evolving average effective optical depth, that the data are best explained by models with reionization completing at $z \\la 2.9$.  This is somewhat later than expected from measurements of the IGM temperature \\citep{schaye00, ricotti00, lidz09-temp}. It is more consistent with hints that the shape of the ionizing background spectrum changes around that time \\citep{songaila98, songaila05}, although that argument has remained controversial \\citep{kim02, aguirre04}.  Late reionization is reasonably consistent with estimates of the total ionizing emissivity of high-$z$ quasars, which produce a couple of ionizing photons per helium atom by this point (e.g., \\citealt{furl08-helium}).  We have focused on variations in the radiation background. \\citet{fardal98} showed that fluctuations induced by large-scale IGM density variations are modest.  While they will slightly modify the values in Tables~\\ref{tab:conf-30} and \\ref{tab:conf-60}, they will not qualitatively modify our conclusions about the high-$\\tau_{\\rm eff}$ measurements.  They can, however, help to reduce the tension at $z \\la 2.3$, where the absolute deviations are relatively small.  Our toy model of reionization can obviously be much improved with realistic simulations of reionization, such as those performed by \\citet{mcquinn09} -- although the large volume required to accumulate enough samples presents a challenge.  The probabilities generated from our simple model only provide a plausibility argument that this event, where large coherent regions with weak radiation fields should be relatively common, can more easily accommodate the observed $\\tau_{\\rm eff}$ distributions.  Qualitatively, we expect similar behavior in more sophisticated models (see, for example, the visualizations in Figs. 4-7 of \\citealt{mcquinn09}).  We therefore have not tried to estimate the evolution of the \\ion{He}{3} fraction with redshift but instead to argue that the data suggest some sort of change in the ionized fraction at $z \\la 2.9$.  Our post-reionization model does require some caveats.  First, it does not include the details of radiative transfer, such as shadowing and spectral filtering.  These effects can clearly cause substantial fluctuations across the absorber population \\citep{maselli05,tittley07}.  However, by averaging over large segments of each spectrum ($\\ga 30 \\Mpc$ across), we should also average over these details.  For example, it is difficult to imagine an absorber large enough to shadow an entire 30 Mpc region from the many sources that illuminate it within distances of a few attenuation lengths.  Second, we have ignored the deterministic clustering of quasars \\citep{shen07}. \\citet{furl09-hefluc} showed that this will be relatively unimportant to the distribution of $\\Gamma$, because the fractional variance in quasar emissivity within one attenuation volume is much larger than that produced by deterministic clustering (thanks to the large Poisson fluctuations and the wide luminosity range of quasars).  \\citet{mcquinn09} also showed that the distribution of \\ion{He}{2} during reionization is essentially driven by the Poisson fluctuations of the quasar population.  However, quasar clustering may be somewhat more important for our application, because clustering \\emph{within} the attenuation volumes cannot really be neglected \\citep{mesinger09} and because the spatial coherence is important for averaging over long skewers.  However, it should still not be a dominant effect, and we expect that our qualitative conclusions will hold.  We have also simplified the problem by taking a single attenuation length for all photons, whereas in reality it increases with frequency.  However, the pair of values we have used spans a factor of two in length, both yielding similar conclusions, so we do not expect this to make a large difference.  If anything, including higher-energy photons will \\emph{strengthen} evidence for late reionization, because larger attenuation lengths decrease the fluctuations.  It may seem surprising that we have drawn such strong conclusions, given the very limited number of lines of sight so far available.  The most conservative conclusion one could draw is that the data are inconsistent with a smoothly evolving ionizing background over the entire range $z \\sim 2$--$3.2$.  Either (1) there must be a sudden increase in the slope of $\\VEV{\\tau_{\\rm eff}}(z)$ or (2) extra fluctuations from reionization are required.  These conclusions are robust because they rely on the difficulty of producing a region  of any substantial length with $\\Gamma \\ll \\bar{\\Gamma}$ in the post-reionization Universe.  Thus, so long as an interval is large enough to include a substantial number of absorbers (and so average over the density fluctuations), the detection of even a single region with $\\tau_{\\rm eff} \\gg \\VEV{\\tau_{\\rm eff}}$ provides strong evidence for reionization.  In any case, the next several years should see a substantial increase in the number of lines of sight, with dozens of promising targets recently detected \\citep{zheng04-sdss, zheng08, syphers09, syphers09b}.  The installation of the Cosmic Origins Spectrograph and repair of the Space Telescope Imaging Spectrograph on the \\emph{Hubble Space Telescope} (HST) have provided two powerful instruments for such exploration; intriguingly, these instruments are best suited for studying the IGM at $z \\ga 3$, precisely where the fluctuations should be most interesting.  \\citet{syphers09, syphers09b} have detected 22 prime candidates for \\ion{He}{2} \\lya forest spectra.  Nearly all lie in the extremely interesting redshift range $z \\sim 3$--$4$, in the heart of the reionization era.  A visual inspection of the modest-resolution spectra available from their surveys (taken with the ACS grism on HST) shows many interesting features both within individual spectra and across the entire population.  Finally, our model shows that even modest-resolution and relatively low signal-to-noise investigations (such as the ACS grism sample) can make substantial progress in understanding \\ion{He}{2} reionization.  We have drawn strong conclusions from the absorption averaged over fairly large intervals ($\\Delta z \\sim 0.03$--$0.1$), which is (at least in principle) a resolution-independent measurement.  Thus, this test can be useful even for quasars that are too faint for the detailed \\ion{He}{2} \\lya forest measurements that have so far held the focus of the community."
    },
    "0910/0910.2466_arXiv.txt": {
        "abstract": "We propose a new heating mechanism in magnetar crusts.  Magnetars' crustal magnetic fields are much stronger than their surface fields; therefore, magnetic pressure partially supports the crust against gravity.  The crust loses magnetic pressure support as the field decays and must compensate by increasing the electron degeneracy pressure; the accompanying increase in the electron Fermi energy induces nonequilibrium, exothermic electron captures.  The total heat released via field-decay electron captures is comparable to the total magnetic energy in the crust.  Thus, field-decay electron captures are an important, if not the primary, mechanism powering magnetars' soft X-ray emission. ",
        "introduction": "Soft gamma repeaters (SGRs) and anomalous X-ray pulsars \\citep[AXPs; for reviews, see][]{WT06,K07,M08} are the two known manifestations of magnetars: neutron stars with ultrastrong exterior magnetic fields $B \\gtrsim 10^{14} \\nsp \\gauss$ and even stronger interior, primarily toroidal fields \\citep{DT92,P92,TD93}.  As of this writing, all magnetars have characteristic ages $\\tage \\lesssim 10^{5} \\nsp \\yr$, and their persistent soft X-ray luminosities $L \\sim 10^{35} \\nsp \\ergspersecond$ \\citep[although some are transients, e.g.,][]{Ietal04}, which are likely of thermal origin, are systematically higher by factors of $>100$ than the power available from rotation \\citep{vPTvdH95}, cooling \\citep[e.g.,][] {PGW06}, or accretion \\citep{HvKK00,Ketal01}.  Taken together, these observations suggest that magnetars' magnetic fields decay and thereby induce heating in the stellar interiors \\citep[e.g.,][]{TD96,HK98,CGP00,PG07,PLMG07}.  Sustaining the persistent thermal emission requires an interior energy reservoir \\begin{equation}\\label{e.persistentE} \\Eper \\gtrsim 3 \\ee{47} \\left ( \\frac{\\tage}{10^{5} \\nsp \\yr} \\right ) \\left (  \\frac{L}{10^{35} \\nsp \\ergspersecond} \\right ) \\nsp \\ergs. \\end{equation} Comparing this to the magnetar's total internal magnetic energy $\\Emag \\sim (B^{2}/8 \\pi) (4 \\pi R^{3}/3)$ gives $B \\gtrsim 1\\ee{15} \\nsp \\gauss$.  However, these values are likely flagrant underestimations:  First, a magnetar emits most of its thermal energy as neutrinos.  Indeed, \\citet{KYPSSG06,KPYC09} argue that, to avoid severe neutrino losses, the heat source powering the thermal emission must be located at or near the outer crust, i.e., within the magnetar's outermost $\\approx 100 \\nsp \\meter$.  If the thermal emission originates from magnetic energy contained only within the outer crust, then $B \\sim 10^{16} \\nsp \\gauss$.  Second, a magnetar's magnetic field powers other phenomena as well, such as giant flares;  the 2004 December 27 giant flare from SGR 1806--20 released $\\sim 10^{46} \\nsp \\ergs$, which implies $B \\sim 10^{16} \\nsp \\gauss$ \\citep[assuming giant flares are recurrent;][]{Hetal05,SDIV05,Tetal05}.  That the magnetic field's strength is much greater in the crust than at the surface implies that a substantial $B$ gradient exists.  Therefore, if the crustal field geometry is primarily nonradial, magnetic pressure partially supports the crust against gravity.  We argue in \\S \\ref{s.fielddecayecaptures} that the crust loses magnetic pressure support as $B$ decays and compensates with an increase in electron degeneracy pressure.  The accompanying increase in the electron Fermi energy induces exothermic electron captures that heat the crust.  We calculate the total energy released via field-decay electron captures in \\S \\ref{s.crustalheating} and compare it to Joule heating in \\S \\ref{s.jouleheatingcomparison}.  We discuss our results and speculate on other consequences of field-decay electron captures in \\S \\ref{s.discussion}. ",
        "conclusions": "\\label{s.discussion}  In this work, we identified a new, strong heat source in magnetar crusts.  Magnetic pressure partially supports the crust against gravity.  The electron degeneracy pressure increases as the crust loses magnetic pressure support during field decay, and the concomitant increase in $\\EF$ induces exothermic electron captures.  The total heat released via field-decay electron captures is of order the crust's total magnetic energy.  Field-decay electron captures are important to neutron stars' magneto-thermal evolution and should be included in future calculations.  At a minimum, field-decay electron captures supplement Joule heating and lessen the interior magnetic field strength required to power magnetar phenomena; however, since only a fraction of the crust's magnetic energy, but essentially all of the nuclear energy, converts to thermal energy, field-decay electron captures may in fact dominate magnetars' persistent soft X-ray emission.  We now outline other potential ramifications of this work:  (1) Although we neglected both the inner crust and outermost envelope for simplicity, field-decay-induced nuclear reactions can occur there as well.  In addition to electron captures in the inner crust, the loss of magnetic pressure support may trigger light-element fusion reactions in the envelope or pycnonuclear reactions in the inner crust, since $\\rho$ increases at a fixed $\\Sigma$ during field decay.  (2) The crust's electrical resistivity decreases with $\\rho$ and increases with $T$ \\citep[e.g.,][]{YU80}, so magnetic field decay via Hall-drift-enhanced Ohmic dissipation \\citep[e.g.,][]{GR92} is most efficient at low densities and high temperatures.  Magnetic pressure support lowers $\\rho$ at a given $\\Sigma$ relative to the nonmagnetic case, and field-decay electron captures heat the crust; both effects accelerate Ohmic dissipation.  (3) Although it is the dominant mechanism, field decay is not required to induce electron captures.  For example, rearrangement of the crustal magnetic field via Hall drift, which is nondissipative, can induce captures:  Altering magnetic pressure gradients alone can either raise or lower the electron pressure, and thereby $\\EF$, of a matter parcel past an electron capture threshold.  Raising $\\EF$ past a capture threshold triggers exothermic electron captures, but lowering $\\EF$ past a threshold does nothing, since exothermic captures are effectively irreversible.  This may further enhance the heating rate, and hence the Ohmic dissipation rate, in the crust.  (4) Asymmetries in the internal magnetic field relative to the stellar rotation axis produce distortions in a young, rapidly rotating magnetar that may generate gravitational waves.  A nascent misalignment between the rotation and magnetic axes is secularly unstable and grows with time \\citep[e.g.,][]{C02,DSS09}.  If the distorting magnetic field is predominantly nonradial and within the crust, the resulting field evolution causes electron captures, which can generate asymmetric density variations that intensify gravitational wave emission \\citep[see][who proposed a similar mechanism for weakly magnetic, accreting neutron stars]{B98,UCB00}.  (5) We neglected the crust's elasticity in the hydrostatic equilibrium equation (\\ref{e.hydroeq}).  Elasticity partially compensates local magnetic stresses and thereby delays electron captures; conversely, sudden pressure increases that offset fracture-induced elasticity losses during catastrophic events like SGR bursts and giant flares may trigger electron captures.   While unlikely to contribute to the burst's emission itself, captures triggered by crust fracture may contribute to the subsequent transient cooling \\citep[e.g.,][]{LET02}."
    },
    "0910/0910.0194_arXiv.txt": {
        "abstract": "Nuclear reactions proceed differently in stellar plasmas than in the laboratory due to the thermal effects in the plasma. On one hand, a target nucleus is bombarded by projectiles distributed in energy with a distribution defined by the plasma temperature. The mostly relevant energies are low by nuclear physics standards and thus require an improved description of low-energy properties, such as optical potentials, required for the calculation of reaction cross sections. Recent studies of low-energy cross sections suggest the necessity of a modification of the proton optical potential. On the other hand, target nuclei are in thermal equilibrium with the plasma and this modifies their reaction cross sections. It is generally expected that this modification is larger for endothermic reactions. We show that there is a large number of exceptions to this rule. ",
        "introduction": "Stellar plasmas are characterized by their composition and temperature. The energy distribution of nuclei according to the temperature is given by a Maxwell-Boltzmann (MB) function. This function peaks at comparatively low energy and thus the astrophysically relevant interaction energies of nuclei are low by nuclear physics standards. This is a challenge for experimental and theoretical nuclear physics, especially when dealing with charged-particle induced reactions. The thermalization of nuclei is very fast, even compared to reaction timescales. This populates excited states and reactions can proceed from the ground state and the excited states, contrary to reactions in the laboratory which only include targets in the ground state. The resulting change in the effective cross section is called stellar enhancement and usually given by the ratio of the stellar cross section to the ground state one, $f=\\sigma^*/\\sigma^\\mathrm{g.s.}$. The stellar enhancement factor $f$ can only be studied theoretically. ",
        "conclusions": ""
    },
    "0910/0910.5064_arXiv.txt": {
        "abstract": "We perform ideal-magnetohydrodynamic axisymmetric simulations of magnetically confined mountains on an accreting neutron star, with masses $\\lesssim 0.12 \\Msun$. We consider two scenarios, in which the mountain sits atop a hard surface or sinks into a soft, fluid base. We find that the ellipticity of the star, due to a mountain grown on a hard surface, approaches $\\sim 2 \\times 10^{-4}$ for accreted masses $\\gtrsim 1.2 \\times 10^{-3} \\Msun$, and that sinking reduces the ellipticity by between 25\\% and 60\\%. The consequences for gravitational radiation from low-mass x-ray binaries are discussed. ",
        "introduction": "The magnetic dipole moment $\\mu$ of a neutron star is observed to diminish in the long term as the star accretes \\citep{Taam-vdHeuvel-1986,vdHeuvel-Bitzaraki-1995}, although \\citet{Wijers-1997} argued that $\\mu$ may also be a function of parameters other than the accreted mass $\\Ma$. The $\\mu$--$\\Ma$ correlation has been ascribed to a number of physical mechanisms \\citep{Melatos-Phinney-2001,Cumming-2005}. First, the magnetic field may be dissipated in the stellar crust by Ohmic decay, accelerated by heating as the accreted plasma impacts upon the star \\citep{Konar-Bhattacharya-1997,Urpin-etal-1998,Brown-Bildsten-1998,Cumming-etal-2004}. Second, magnetic flux tubes may be dragged from the superconducting core by the outward motion of superfluid vortices, as the star spins down \\citep{Srinivasan-etal-1990,Ruderman-etal-1998,Konar-Bhattacharya-1999,Konen-Geppert-2001}. Third, the magnetic field may be screened by accretion-induced currents within the crust \\citep{BisnovatyiKogan-Komberg-1974,Blondin-Freese-1986,Lovelace-etal-2005}. In particular, the field may be \\emph{buried} under a mountain of accreted plasma channelled onto the magnetic poles. When $\\Ma$ is large enough, the mountain spreads laterally, transporting the polar magnetic flux towards the equator \\citep{Hameury-etal-1983,Romani-1990,Brown-Bildsten-1998,Cumming-etal-01,Melatos-Phinney-2001, Choudhuri-Konar-2002,PM04,Zhang-Kojima-2006,PM07,VM08,VM09R}.  \\citet{PM04} computed the unique sequence of self-consistent, ideal-magneto\\-hydro\\-dynamic (ideal-MHD) equilibria that describes the formation of a polar mountain by magnetic burial as a function of $\\Ma$. They found that the accreted mountain is confined by the equatorially compressed magnetic field, which was unaccounted for in previous calculations, and that $10^{-5}\\Msun$ must be accreted to lower $\\mu$ by 10\\%. Surprisingly, mountains are stable with respect to axisymmetric ideal-MHD perturbations; they oscillate globally in a superposition of acoustic and Alfv\\'en modes but remain intact due to magnetic line-tying at the stellar surface \\citep{PM07}. The same equilibria are susceptible to nonaxisymmetric, Parker-like instabilities (specifically the gravitationally driven, undular sub-mode), but the instability preserves a polar mountain when it saturates, despite reducing the mass ellipticity by $\\sim 30\\%$ \\citep{VM08}. Recently, \\citet{VM09R} considered resistive effects. They found that the mountain does not relax appreciably for realistic resistivities over the lifetime of a low- or high-mass X-ray binary, either by global diffusion, resistive g-mode instabilities, or reconnection in the equatorial magnetic belt. The Hall drift, which exerts a destabilising influence in isolated neutron stars \\citep[see, e.g.,][]{Rheinhardt-Geppert-2002}, is unlikely to be important in accreting neutron stars due to crustal impurities \\citep{Cumming-etal-2004,Cumming-2005}.  The investigations outlined in the previous paragraph suffer from two limitations. First, the mountain is assumed to rest upon a rigid surface. Under this assumption, the accreting plasma cannot sink into the stellar crust. This is unrealistic. During magnetic burial, frozen-in magnetic flux is redistributed slowly within the neutron star by the accreted plasma, as it sinks beneath the surface and spreads laterally. \\citet{Choudhuri-Konar-2002} showed that the time-scale and end state of burial are tied to these slow interior motions. Second, the accreted plasma is assumed to satisfy an isothermal equation of state. This is an accurate model only for neutron stars with low accretion rates $\\Madot \\lesssim 10^{-10} \\Msun \\mathrm{yr}^{-1}$; the thermodynamics of neutron stars accreting near the Eddington limit ($\\sim 10^{-8} \\Msun \\mathrm{yr}^{-1}$) is more complicated, with a depth-dependent adiabatic index \\citep{Brown-Bildsten-1998,Brown-2000}. The equation of state affects the growth rate of Parker-like instabilities \\citep{Kosinski-Hanasz-2006}.  In this paper, we seek to overcome the first limitation. In section~\\ref{sec:method}, we present a new method of computationally simulating the growth of a magnetic mountain with $\\Ma \\lesssim 0.1 \\Msun$. In section~\\ref{sec:compare}, we compare the structure of mountains grown on hard and soft surfaces to evaluate the role of sinking. In section~\\ref{sec:ellip}, the resulting mass quadrupole moment is evaluated as a function of $\\Ma$ for hard and soft surfaces. A comparison with the results of \\citet{Choudhuri-Konar-2002}, and the implications for gravitational wave emission from rapidly rotating accretors (e.g. low-mass X-ray binaries), are discussed in section~\\ref{sec:discuss}. ",
        "conclusions": "\\label{sec:discuss}  \\begin{table} \\centering \\caption{\\label{tbl:parameters} List of important physical parameters of accreting neutron stars (top part), and a summary of the results of the simulations presented in this paper (bottom part). } \\begin{tabular}{@{}lll} \\hline  Quantity & Value/Range & Reference \\\\  \\hline  accreted mass & $10^{-4}$ -- $0.8 \\Msun$ & 1, 2, 4 \\\\  accretion timescale & $10^{4}$ -- $10^{6}~\\mathrm{yr}$ & 1, 7 \\\\  density of crust & $10^{9}$ -- $10^{14}~\\mathrm{g~cm}^{-3}$ & 5, 10 \\\\  depth of crust & $\\sim 1000~\\mathrm{m}$ & 5, 10 \\\\  initial magnetic field & $10^{12}$ -- $10^{13}~\\mathrm{G}$ & 3, 8, 9 \\\\  temperature & $10^{8}$ -- $10^{9}~\\mathrm{K}$ & 6, 10 \\\\  \\hline  ellipticity & $5 \\times 10^{-5}$ -- $2 \\times 10^{-4}$ & 11 \\\\  effect of sinking & $\\epsilon$ reduced by 25 -- 60 \\% & 11 \\\\  \\hline \\end{tabular} \\begin{flushleft} References: 1.~\\citet{Taam-vdHeuvel-1986}; 2.~\\citet{vdHeuvel-Bitzaraki-1995}; 3.~\\citet{Hartman-etal-1997}; 4.~\\citet{Wijers-1997}; 5.~\\citet{Brown-Bildsten-1998}; 6.~\\citet{Brown-2000}; 7.~\\citet{Cumming-etal-01}; 8.~\\citet{Arzoumanian-etal-2002}; 9.~\\citet{FaucherGiguere-Kaspi-2006}; 10.~\\citet{Chamel-Haensel-2008}; 11.~This work. \\end{flushleft} \\end{table}  In this paper, we simulate the growth of magnetically confined mountains on an accreting neutron star, with realistic masses $\\Ma \\lesssim 0.12 \\Msun$, under the two scenarios where the mountain sits on a hard surface and sinks into a soft, fluid base. In the latter scenario, we confirm that the final equilibriun state is independent of the altitude where matter is injected. We find that the ellipticity of a hard-surface mountain does not increase appreciably for $\\Ma \\gtrsim \\Mii\\Mc$, saturating at $\\sim 2 \\times 10^{-4}$, whereas the ellipticity of a soft-surface mountain continues to increase from $\\Ma = \\Mii\\Mc$ to $\\Miiii\\Mc$. Sinking reduces the ellipticity by up to 60\\% relative to the hard-surface value.  \\citet{Choudhuri-Konar-2002} developed a kinematic model of accretion, which treats sinking in a different (but complementary) way to this paper. An axisymmetric magnetic field is evolved under the influence of a \\emph{prescribed} velocity field, which models the flow of accreted matter from pole to equator, where it submerges and moves towards the core (see their Figure~1). Ohmic diffusion is included, but, for a subset of the results (where the resistivity $\\eta = 0.01$), it is negligible, permitting a direct comparison with this paper.  Figure~5 of \\citet{Choudhuri-Konar-2002} shows the evolved configuration of an initially dipolar field which permeates the entire star. We compare to Figures~\\subref*{fig:flux-sink-1-t21} and~\\subref*{fig:flux-sink-1-t21} of this paper. In both models, the magnetic field is distorted significantly by accreted matter spreading towards the equator ($r_\\mathrm{m} < r < r_\\mathrm{s}$ in \\citeauthor{Choudhuri-Konar-2002}; the entire simulation in this paper). In \\citeauthor{Choudhuri-Konar-2002}'s work, the magnetic field is completely submerged beneath the surface and confined to the zone where the submerged accreted matter flows back towards the pole. In this paper, magnetic field lines still penetrate the surface, implying less effective screening. Within the core, \\citeauthor{Choudhuri-Konar-2002}'s magnetic field remains relatively undisturbed. Magnetic line-tying is not enforced, but the prescribed radial flow within the core naturally restricts the sideways displacement of the magnetic field there. If there were sideways motion of the matter within the core, it would modify the degree of magnetic screening, but neither our simulations nor the results of \\citeauthor{Choudhuri-Konar-2002} show evidence for such motion. Extending our simulations deeper into the star to include the core and explore this possibility properly would be a technical challenge; for instance, we would need to incorporate a more realistic equation of state and track even more disparate equilibrium time-scales.  To explain the narrow range in the rotation frequencies of low-mass x-ray binaries \\citep{Chakrabarty-etal-2003}, it is proposed that the stars radiate angular moment in gravitational waves at a rate which balances the accretion torque \\citep{Wagoner-1984,Bildsten-1998}. Magnetic mountains are one of a number of physical mechanisms proposed for the associated permanent quadrupole; see \\citet{VM09GW} and references therein. The relationship between $\\epsilon$ and the rotation frequency $f$ predicted by torque balance is $f \\propto \\epsilon^{-2/5}$. Thus, the 25\\% to 60\\% reduction in $\\epsilon$ due to sinking calculated in this paper increases $f$ by 12\\% to 44\\%, all other things being equal. This goes some way towards bringing magnetic mountain ellipticities down to a level consistent with the data, but there is still a long way to go. Observations to date have found $45 \\text{ Hz} < f < 620 \\text{ Hz}$ for burst oscillation sources and $182 \\text{ Hz} < f < 598 \\text{ Hz}$ for accreting millisecond pulsars, implying $6.6 \\times 10^{-9} \\lesssim \\epsilon \\lesssim 4.6 \\times 10^{-6}$ and $7.2 < 10^{-9} \\lesssim \\epsilon \\lesssim 1.4 \\times 10^{-7}$ respectively. Conversely, the ellipticities of sunk mountains calculated in this paper, $3.5 \\times 10^{-5} \\lesssim \\epsilon \\lesssim 1.5 \\times 10^{-4}$, imply $11 \\text{ Hz} \\lesssim f \\lesssim 20 \\text{ Hz}$. Clearly, other relaxation mechanisms, like Ohmic diffusion, must also be playing an important role in reducing $\\epsilon$, as the observed $f$ require.  The reduction in $\\epsilon$ by sinking also reduces the gravitational wave strain \\citep[e.g.][]{LSC-CW-Fstat-S2}, $h \\propto \\epsilon f^2$, by 6\\% to 17\\%. This is unlikely, by itself, to rule out the detection of gravitational waves from low-mass x-ray binaries by ground-based interferometric detectors; assuming the signal can be coherently integrated, the loss in $h$ can be compensated for by an increase in the observation time $\\propto h^{-2}$ of 13\\% to 45\\%. Other difficulties associated with the detection of gravitational waves from low-mass x-ray binaries, such as poorly known orbital parameters and accretion-induced phase wandering \\citep{Watts-etal-2008}, are likely to be more important."
    },
    "0910/0910.5913_arXiv.txt": {
        "abstract": "{Until now, most members of the Ursa Major (UMa) group of stars have been identified by means of kinematic criteria. However, in many cases kinematic criteria alone are insufficient to ascertain, whether an individual star is really a member of this group. Since photometric criteria are ineffective in the case of cool dwarf members, one must use spectroscopic criteria. Nevertheless, resulting membership criteria are inconclusive.} {We reanalyse spectroscopic properties of cool UMa group dwarfs. In particular, we study the distribution of iron abundance, the strength of the Li\\,I absorption at $6708\\,\\mathrm{\\AA}$ and the Li abundance, and the infilling of the core of the H$\\alpha$ line.} {Twenty-five cool and northern bona-fide members are carefully selected from the literature. Homogeneously measured stellar parameters and iron abundances are given for all Sun-like stars selected, based on spectra of high resolution and high signal-to-noise ratio. In addition, we measure the Li equivalent width and abundance as well as the relative intensity of the H$\\alpha$ core and the corresponding chromospheric flux.} {The studied stars infer an average Ursa Major group iron abundance of $-0.03\\pm0.05\\,$dex, which is higher by about $0.06\\,$dex than determined elsewhere. The Li abundance derived of Ursa Major group dwarf stars is higher than in the Hyades at effective temperatures cooler than the Sun, but lower than in the younger Pleiades, a result which is independent of the exact value of the effective temperature adopted. The Sun-like and cooler dwarfs also display chromospheric infilling of the H$\\alpha$ core. We present spectroscopic criteria that may be used to exclude non-members.} {} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.2702_arXiv.txt": {
        "abstract": "We explore the possibility of a local origin for ultra high energy cosmic rays (UHECRs). Using the catalogue of Karachentsev et al. including nearby galaxies with distances less than 10Mpc (Local Volume), we search for a correlation with the sample of UHECR events released so far by the Pierre Auger collaboration. The counterpart sample selection is performed with variable distance and luminosity cuts which extract the most likely sources in the catalogue. The probability of chance correlation after penalizing for scans is 0.96\\%, which corresponds to a correlation signal of $2.6\\sigma$. We find that the parameters that maximize the signal are $\\psi=3.0^{\\circ}$, $D_{\\rmn{max}}=4$Mpc and $M_{\\rmn{B}}=-15$ for the maximum angular separation between cosmic rays and galaxy sources, maximum distance to the source, and sources brighter than $B$-band absolute magnitude respectively.  This implies a preference for the UHECRs arrival directions to be correlated with the nearest and most luminous galaxies in the Local Volume, while the angular distance between the cosmic ray events and their possible sources is similar to that found by The Pierre Auger Collaboration using active galactic nuclei (AGNs) within 70-100Mpc instead of local galaxies. We note that nearby galaxies with $D<10$Mpc show a similar correlation with UHECRs as compared to well-known particle accelerators such as AGNs, although less than 20\\% of cosmic ray events are correlated to a source in our study. However, the observational evidence for mixed composition in the high-energy end of the cosmic ray spectrum supports the possibility of a local origin for UHECRs, as CNO nuclei can travel only few Mpc without strong attenuation by the GZK effect, whereas the observed suppression in the energy spectrum would require more distant sources in the case of pure proton composition interacting with the CMB. ",
        "introduction": "Ultra-high energy cosmic rays (UHECRs) may become a complementary probe of some astrophysical objects in addition to observations in multiple wavelengths. In fact, the detection of these particles with enough statistics would represent the awake of the development of multi-messenger astrophysics. The Pierre Auger observatory in Malarg\\\"ue, Argentina (see \\citealt{2004NIMPA.523...50A}) has devoted a large effort in the construction of an unprecedented array of water {\\v C}erenkov tanks covering an area of $\\sim$3,000 km$^2$, together with four fluorescence telescopes which allow for increased accuracy in energy measurements. This effort has proven to be fruitful as the first relevant results were obtained even before the completion of the entire experiment array. In particular, the unprecedented statistics on the detection of ultra-high energy cosmic rays above 10EeV=$10^{19}eV\\simeq 1J$ with arrival directions measured with an accuracy better than $1^{\\circ}$ has allowed to search for possible astrophysical sources of these particles with higher reliability than previous experiments (e.g. \\citealt{2004APh....21..359F}).  By the end of 2007, the Pierre Auger Collaboration concluded that the arrival directions of UHECRs above 55EeV correlate with the positions of nearby active galactic nuclei (AGN), or other objects which trace, in the same way as AGNs, the Large Scale Structure of the Universe (\\citealt{2007Sci...318..938T}, \\citealt{2008APh....29..188T}). Statistics of these cosmic ray events has doubled since then (from 27 to 58 events), and yet the correlation with these objects has not strengthened as compared to previous estimations \\citep{2009arXiv0906.2347T}. This supports the fact that this scenario is not the only one which can reproduce the observed data, although increasing statistics of UHECR events may be able to finally disentangle their real origin. Expectations in the cosmic-ray community are optimistic, and as stated in \\citet{2009astro2010S.225O}, this issue may find a solution in the next decade.  Indeed, there has been a plethora of studies in the literature trying to find out a correlation of the UHECRs with other candidate sources such as particular objects like galaxy clusters, GRBs or other potential hosts of energetic phenomena (e.g. \\citealt{2005astro.ph..7679P}, \\citealt{2008PhRvD..78b3005M}, \\citealt{2008MNRAS.390L..88G}) and structures like the Supergalactic Plane \\citep{2008arXiv0805.1746S}. Some of these works did report a similar significance as compared to the Auger result on the correlation with the angular positions of nearby AGNs. These similar values of the probability of chance correlation failed to discriminate with enough robustness between different cosmic-ray acceleration sites. In order to get stronger conclusions from this scarce set of UHECR events, there have been several efforts on the development of new methods to estimate the chance probability of correlation in the case of isotropic flux with potential sources (\\citealt{2009JCAP...07..023A}, \\citealt{2009arXiv0906.2347T}). Despite all these efforts, there is no preferred candidate source yet.  Therefore, any complementary information from other measurements, like the shape of the energy spectrum, or the composition of these cosmic rays, plays now a decisive role. For example, interactions with CMB photons via the Greisen-Zatsepin-Kuzmin effect (\\citealt{1966PhRvL..16..748G}, \\citealt{1966JETPL...4...78Z}) are expected to decimate the statistics of the UHECRs arriving to the Earth as a function of their composition and travel distance (e.g. \\citealt{2005NuPhS.138..465C}, \\citealt{2009astro2010S.225O}). Recent claims of a detection of this GZK cutoff (\\citealt{2008PhRvL.100j1101A}, \\citealt{2008PhRvL.101f1101A}) raises the question of how far away these cosmic accelerators must be located in order to present the observed suppression. Besides, simple considerations regarding the coincidence of the GZK and UHECR clustering energies point to the fact that the composition of UHECRs must deviate from pure proton composition and tend to be CNO--like \\citep{2008JPhCS.120f2006D}, which is favoured by observations (\\citealt{2009arXiv0901.3389B}, \\citealt{2009arXiv0902.3787U}, \\citealt{2009arXiv0906.2319T}). As CNO nuclei can travel shorter distances as compared to protons, this sets a strong limit on the distance of UHECR sources, without suffering from the large deflection angles by the galactic magnetic field in the case of iron nuclei \\citep{2009arXiv0909.1532T}.  In this paper we address the possibility of a local origin for the Auger UHECRs. We use the only data published so far in \\citet{2008APh....29..188T}, which comprises 27 UHECR events detected prior to 1 September 2007, with energies above 57EeV (which has been revised to 55EeV due to new calibration procedures). This energy scale turns out to be however very interesting in the search for these comic ray sources, as \\citet{2009arXiv0901.4699M} showed that the clustering above this energy is statistically significant. Most of the recent studies, including \\citet{2008APh....29..188T}, have focused so far on sources located as far as $~75-100$Mpc, which corresponds to the GZK horizon for protons. Here we assess if there is any correlation signal with nearby objects with distance $D<10$Mpc (Local Volume) which is roughly the GZK horizon for CNO nuclei. In order to explore this hypothesis, we take advantage of the Local Volume catalogue of galaxies of \\citet{2004AJ....127.2031K}, which is about 70--80\\% complete up to 8Mpc. This allows us to evaluate the likelihood of galaxies in the Local Volume as UHECR accelerator sites.  This paper is organized as follows: in Section~\\ref{sec:methodology} we describe the methodology used in our analysis and the catalogue of nearby galaxies in the Local Volume. In Section~\\ref{sec:results} we present the results of the chance probability of the correlation of UHECRs with these objects. We discuss our results and conclude in Section~\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion} We have explored the possibility of the correlation of nearby galaxies from the catalogue of \\citet{2004AJ....127.2031K} with ultra high energy cosmic ray events from the Pierre Auger Collaboration. This analysis is highly motivated in the light that a mixed CNO (as opposed to pure proton) composition of UHECRs, which is the preferred scenario indicated by current data, could not survive the GZK cutoff unless the sources are located in the Local Volume. Therefore, any correlation with these nearby sources would explain both composition and correlation at the same time. The analysis of the data concludes a $2.6\\sigma$ correlation between both samples. There are 5 out of 27 correlating events, while 0.5 should be expected on average in the case of isotropic flux. Due to the small number of correlating events, we cannot specify the scan parameters a priori, thus making our results sensitive to the negative impact of a posteriori anisotropy searches. The probability of chance correlation due to isotropic flux after penalizing for scans is $P_{\\rmn{chance}}=0.96\\times 10^{-2}$, rejecting the hypothesis of isotropy of arrival directions of UHECRs at 99\\% confidence level. For comparison, the same result before penalizing for scans is $\\sim 1\\times 10^{-4}$ which is similar to the $\\sim 2\\times 10^{-4}$ result (with scan parameters set a priori) from the correlation found by \\citet{2008APh....29..188T} with the AGNs in the V{\\'e}ron-Cetty catalogue, and stronger than the revised value of $\\sim 6\\times 10^{-3}$ when they include the new UHECR events from 1 September 2007 to 31 March 2009 \\citep{2009arXiv0906.2347T}. This is an indication that there is no clear source of the UHECRs, and finding a lower value of the probability of chance correlation does not guarantee that this is actually the preferred source as compared to other sources which present similar values of $P_{\\rmn{chance}}$. This was already noticed even in early papers such as \\citet{2001JETPL..74..445T}, and in any case it seems unlikely to get a $>3.3\\sigma$ detection \\citep{2008JCAP...05..006K}. Interestingly, most of the objects in the catalogue of \\citet{2004AJ....127.2031K} which correlate with Auger events, are already included in the AGN catalogue, which has around 30 objects with $D<10$Mpc. This may indicate that the relation between nearby galaxies and UHECRs found in our work could be inherited from the \\citet{2008APh....29..188T} result. In any case, the so-called AGN hypothesis has proven to be controversial (\\citealt{2008JETPL..87..461G}, \\citealt{2009ApJ...693.1261M}), stressed by the deficit of events from the direction of the Virgo supercluster. On the other hand, the correlation with the distribution of luminous matter seems to be well established \\citep{2008JCAP...05..006K}, without making any reference to any particular source which can accelerate particles up to these energies.  The small number of correlating events with nearby galaxies might also point to the possibility of the absence of UHECR accelerators in our local environment, although there is no definitive answer yet. Surprisingly, most of the objects in the \\citet{2004AJ....127.2031K} catalogue which correlate with the UHECR sample (with the exception of NGC300) are located in a small region in the sky, with Supergalactic coordinates $150^{\\circ}<SGL<165^{\\circ}$ and $-10^{\\circ}<SGB<-5^{\\circ}$. This region is near the Supergalactic plane, pointing to preferential arrival directions from nearby galaxies located in a high density region. Future measurements and hopefully the next data release with a total of 58 events as announced in \\citet{2009arXiv0906.2347T} will help to clarify the origin of nearby and distant UHECRs, which will allow to study more in detail the acceleration physics behind them.\\\\  The authors acknowledge support from the Spanish MEC under grant PNAYA 2005-07789. A.J.C. thanks support of the MEC through Spanish grant FPU AP2005-1826. We especially want to thank Vasiliki Pavlidou from CalTech Astronomy Department for useful comments on this paper. This work made extensive use of the x4600 machine at CICA (Spain). We thank Claudio Arjona and Ana Silva at CICA for making this machine available to us."
    },
    "0910/0910.2228_arXiv.txt": {
        "abstract": "Violent relaxation -- the protocluster dynamical response to the expulsion of its residual star forming gas -- is a short albeit crucial episode in the evolution of star clusters and star cluster systems.  Because it is heavily driven by cluster formation and environmental conditions, it is a potentially highly rewarding phase in terms of probing star formation and galaxy evolution.  In this contribution I review how cluster formation and environmental conditions affect the shape of the young cluster mass function and the relation between the present star formation rate of galaxies and the mass of their young most massive cluster. ",
        "introduction": "Compact gas-embedded clusters are observed to be the formation sites of a significant fraction of stars in the local Universe.  Following star formation, feedback processes -- ranging from protostellar outflows, photoionization-driven gas overpressure, stellar winds to, eventually, supernovae -- expel the gas not yet converted into stars.  That is, the efficiency with which a cluster parent core turns its gas into stars is always smaller than 100 per cent.  Not only does gas expulsion terminate star formation, it also leaves the newly-formed cluster out-of-equilibrium, thereby forcing it into violent relaxation.  In an attempt to regain a state of equilibrium, the cluster expands, loses a fraction of its stars or, even, undergo complete disruption.  These are the {\\it cluster infant weight-loss} and {\\it cluster infant mortality}, respectively.  The dynamical response of clusters to gas expulsion implies that the cluster age distribution, the cluster mass function, the ratio between the mass in clusters and the mass in stars, as well as the distribution function of cluster half-mass radii all constitute promising diagnostic tools of star cluster formation conditions.  \\cite[Baumgardt \\& Kroupa (2007)]{BaumgardtKroupa07} have built and made available to the community an $N$-body model grid describing the temporal evolution of cluster mass-loss and cluster spatial expansion for a wide range of local star formation efficiencies (SFE; i.e. the mass fraction of dense molecular gas that cluster-forming cores turn into stars), gas-expulsion time-scales and external tidal field impacts.  To date, this is the most comprehensive grid of model clusters through violent relaxation.  Combined to the ever increasing quality of cluster observational data-sets (including the much needed cluster half-mass radius, see e.g. \\cite[Scheepmaker et al.\\,2007]{Scheepmaker+07}), the time is now ripe to decipher observed properties of young star clusters so as to probe their formation and environmental conditions.  The outline is as follows.  Following a brief reminder of violent relaxation chief parameters in Section \\ref{sec:parmentier:para}, I highlight how formation and environmental conditions affect the shape of the cluster mass function in Section \\ref{sec:parmentier:mf}.  I also present some preliminary results about ongoing work regarding the relation between the luminosity of the brightest young cluster and the present star formation rate (SFR) of spiral galaxies, namely, the $SFR-M_V^{brightest}$ relation (Section \\ref{sec:parmentier:brightest}). ",
        "conclusions": ""
    },
    "0910/0910.0707_arXiv.txt": {
        "abstract": "{ We report on a detailed abundance analysis of two strongly r-process enhanced, very metal-poor stars newly discovered in the HERES project, \\object{CS~29491$-$069} ($\\mathrm{[Fe/H]}=-2.51$, $\\mathrm{[r/Fe]}=+1.1$) and \\object{HE~1219$-$0312} ($\\mathrm{[Fe/H]}=-2.96$, $\\mathrm{[r/Fe]}=+1.5$). The analysis is based on high-quality VLT/UVES spectra and MARCS model atmospheres. We detect lines of 15 heavy elements in the spectrum of \\object{CS~29491$-$069}, and 18 in \\object{HE~1219$-$0312}; in both cases including the \\ion{Th}{II} 4019\\,{\\AA} line. The heavy-element abundance patterns of these two stars are mostly well-matched to scaled solar residual abundances not formed by the s-process. We also compare the observed pattern with recent high-entropy wind (HEW) calculations, which assume core-collapse supernovae of massive stars as the astrophysical environment for the r-process, and find good agreement for most lanthanides. The abundance ratios of the lighter elements strontium, yttrium, and zirconium, which are presumably not formed by the main r-process, are reproduced well by the model. Radioactive dating for \\object{CS~29491$-$069} with the observed thorium and rare-earth element abundance pairs results in an average age of $9.5$\\,Gyr, when based on solar r-process residuals, and $17.6$\\,Gyr, when using HEW model predictions. Chronometry seems to fail in the case of \\object{HE~1219$-$0312}, resulting in a negative age due to its high thorium abundance. \\object{HE~1219$-$0312} could therefore exhibit an overabundance of the heaviest elements, which is sometimes called an ``actinide boost''. } ",
        "introduction": "\\label{sec:intro}  The ages of the oldest stars in the Galaxy provide an important constraint for the time of the onset of stellar nucleosynthesis, with further implications for galaxy formation and evolution. Among the various methods used for age determinations of old stars, nucleochronometry based on long-lived radioactive isotopes, e.g. $^{232}$Th (half-life $\\tau_{1/2}=14.05$\\,Gyr) or $^{238}$U ($\\tau_{1/2} = 4.468$\\,Gyr), has attracted a lot of attention during the last decade. Radioactive decay ages can be derived by comparing the observed abundances of these elements, relative to a stable r-process element, to the production ratio expected from theoretical r-process yields.  The interest in this age determination method was increased by discoveries of very metal-poor stars that are strongly enhanced in r-process elements (i.e., stars having $\\mbox{[Eu/Fe]} > +1.0$ and $\\mbox{[Ba/Eu]} < 0$; hereafter r-II stars, following \\citealt{Beers/Christlieb:2005}), for which the measurement of Th and Eu abundances was possible, and whose stellar matter presumably has experienced only a single nucleosynthesis event. The Th/Eu chronometer has been used for determining the age of, e.g., the progenitor of the r-II star \\object{CS~22892$-$052} \\citep{Snedenetal:1996,Cowanetal:1997,Cowanetal:1999,Snedenetal:2003}. Uranium was first detected in a very metal-poor star, \\object{CS~31082$-$001}, by \\citet{Cayreletal:2001}. \\citet{Cowanetal:2002} tentatively detected U in \\object{BD\\,$+17^{\\circ}3248$}, and recently \\citet{Frebeletal:2007b} have clearly seen U in \\object{HE~1523$-$0901}.  However, it is observationally very challenging to measure the uranium abundance of metal-poor stars, because even in cool, strongly r-process-enhanced, metal-poor giants the strongest U line observable with ground-based telescopes, at $\\lambda = 3859.571$\\,{\\AA}, has an equivalent width of only a few m{\\AA}. Therefore, very high signal-to-noise ratio ($S/N$) spectra are required. Furthermore, the line is detectable only in metal-poor stars without overabundances of carbon or nitrogen, since it is blended with a molecular CN feature. As a result, a U abundance can be measured only in about 1 out of $10^6$ halo stars, while the fraction of halo stars for which abundances of Th and Eu (or similar r-process elements) can be determined is about one order of magnitude larger.  However, a significant complication is that the Th/Eu chronometer seems to fail in some r-process enhanced metal-poor stars, resulting in negative age estimates. For example, \\citet{Hilletal:2002} report $\\log\\left(\\mbox{Th/Eu}\\right) = -0.22$ for \\object{CS~31082$-$001}, and \\citet{Hondaetal:2004b} derive $\\log\\left(\\mbox{Th/Eu}\\right) = -0.10$ for \\object{CS~30306$-$132}; the abundances of the heaviest r-process elements (third-peak-elements, actinides) seem enhanced in these stars with respect to the lanthanide Eu. By comparison, \\citet{Snedenetal:2003} measure $\\log\\left(\\mbox{Th/Eu}\\right)=-0.62$ in \\object{CS~22892$-$052}, and, using a production ratio of $\\log\\left(\\mbox{Th/Eu}\\right)_0 = -0.35$, derive an age of $12.8\\pm3$\\,Gyr. These results clearly cast doubts on the reliability of the Th/Eu chronometer pair; the observed relative abundances of intermediate and some heavy r-process elements diverge. For stars where uranium cannot be detected, third-peak elements such as osmium (if available) seem better partners for age determinations, as theoretical predictions of their corresponding r-process yields are more robust than those for pairs of heavy and intermediate mass elements. However, the dominant transitions of third-peak elements, which are detected as neutral species, lie in the UV. They are therefore hard to employ for reliable abundance analyses. Moreover, abundance ratios with rare-earth elements are more sensitive to uncertainties in the model atmosphere, since these are mostly measured using ionized species \\citep[see also][]{Kratzetal:2007}.  Different astrophysical sites for the r-process have been suggested in the past, including core-collapse supernovae of massive stars, neutron-star mergers and more exotic candidates, but none of these scenarios have so far been proven. The para\\-meters of r-process models therefore had to be defined site-independently; they were nevertheless able to reproduce observed abundance patterns of heavy neutron-capture elements in the Sun and metal-poor stars (see, e.g., papers by \\cite{Kratzetal:1993}, \\cite{Pfeifferetal:2001}, \\citet{Wanajoetal:2002} and \\citet{Kratzetal:2007}). Recent studies showed that lighter elements such as strontium, yttrium and zirconium exhibit more complex behavior, which cannot be explained in a simple r-process scenario. We compare our abundance measurements to the nucleosynthetic yields of a new generation of high-entropy wind (HEW) models which include additional charged-particle processes \\citep{Farouqi:2005,Farouqietal:2005,Farouqietal:2008a,Farouqietal:2008b}.  This paper continues our series on the Hamburg/ESO R-process-Enhanced Star survey (HERES). A detailed description of the project and its aims can be found in \\citet[][ hereafter Paper~I]{HERESpaperI}; methods of automated abundance analysis of high-resolution ``snapshot'' spectra have been described in \\citet[][ hereafter Paper~II]{HERESpaperII}. In this paper we report on detailed abundance analyses of the r-II stars \\object{CS~29491$-$069} and \\object{HE~1219$-$0312} (Sect.~\\ref{sec:AbundanceAnalysis}) based on high-quality VLT/UVES spectra (for details see Sect.~\\ref{sec:ObservationReduction}) and MARCS model atmospheres. Our results are presented in Sect.~\\ref{sec:Results} and discussed in Sect.~\\ref{sec:DiscussionConclusions}. ",
        "conclusions": "\\label{sec:DiscussionConclusions}  Currently there are 12 r-II stars reported in the literature: \\object{CS~22892$-$052} \\citep{Snedenetal:1996}, \\object{CS~31082$-$001} \\citep{Cayreletal:2001,Hilletal:2002}, \\object{CS~29497$-$004} \\citep{HERESpaperI}, \\object{CS~22183$-$031} \\citep{Hondaetal:2004b}, \\object{HE~1523$-$0901} \\citep{Frebeletal:2007b}, and seven additional stars published in \\citet{HERESpaperII}. Among these stars, published abundance analyses based on high-resolution spectroscopy of sufficient quality to detect the \\ion{Th}{II} $4019$\\,{\\AA} line were previously available for only four stars: \\object{CS~22892$-$052}, \\object{CS~31082$-$001}, \\object{CS~29497$-$004}, and \\object{HE~1523$-$0901}. With \\object{CS~29491$-$069} and \\object{HE~1219$-$0312}, we add two stars to this well-studied sample, for a total of six.  The relative abundances of most neutron-capture elements that we analyzed are consistent with their corresponding solar residuals, besides a slight upward trend with the atomic number. However, the progenitor gas cloud of \\object{HE~1219$-$0312} may have experienced a particularly strong enrichment with the heaviest elements. This leads to a failure of the commonly used Th/Eu chronometer, along with most other element pairs, by resulting in a negative radioactive decay age. Selective enhancement of elements in the third r-process peak and beyond requires r-process models with high neutron densities (and entropies in a HEW scenario). The plausibility of such a physical environment therefore needs to be further investigated. Measuring abundances of lead and other third-peak elements, which should exhibit similar enrichment as thorium, could contribute to solving this problem. The low Pb abundance of \\object{CS~31082$-$001} found by \\citet{Plezetal:2004}, however, seems to point towards a more complicated picture. The slight trend in the individual decay age estimates could indicate an alternative explanation of strong actinide enrichment by a deficiency of lighter elements, but the scatter is too large to provide adequate evidence.  Note that \\citet{Hondaetal:2004b} reported an increased Th/Eu ratio of $\\log\\left(\\mathrm{Th/Eu}\\right)=-0.10$ for \\object{CS~30306$-$132}, an r-I giant star with $\\mathrm{[Fe/H]}=-2.42$ and $\\mathrm{[Eu/Fe]}=+0.85$. The case appears similar to that of \\object{HE~1219$-$0312}, with \\object{CS~30306$-$132} having a relative thorium overabundance. Likewise, there are no measurements of other third-peak elements or actinides to further investigate a process that causes strong actinide enrichment.  \\object{CS~29491$-$069} seems to have a significantly smaller thorium abundance with respect to the overall enrichment with heavy elements; an actinide overabundance is therefore unlikely. Radioactive dating based on solar r-process residuals results in an average age of $9.5$\\,Gyr, and $17.6$\\,Gyr for the HEW predictions. The Th/Eu pair seems to yield a much younger age, caused by the low europium abundance. The large scatter in decay ages found for different element pairs confirms that stellar chronometry needs to be based on more than one abundance ratio. The accuracy of absolute ages that are determined from theoretical predictions is still limited, owing to uncertainties in the current nucleosynthesis models and to the unknown astrophysical site of the r-process, apart from the inevitable measurement errors.  We also compare our abundance measurements with the yields of recent dynamical network calculations in the framework of a high-entropy-wind (HEW) scenario. Most heavy elements beyond the second r-process peak show good agreement with the predictions. The model matches our observed abundance ratios of strontium, yttrium and zirconium; however, their absolute values are significantly over-predicted. Further supported by the mismatch of palladium in \\object{HE~1219$-$0312}, this may be interpreted as an indication of an incomplete ejection or fallback scenario for lighter elements, or as contributions from different types of SN II."
    },
    "0910/0910.1701_arXiv.txt": {
        "abstract": "{To understand the mechanisms that trigger solar flares, we require models describing and quantifying observable responses to the original energy release process, since the coronal energy release site itself cannot be resolved with current technical equipment. Testing the usefulness of a particular model requires the comparison of its predictions with flare observations.} {To test the standard flare model (CSHKP-model), we measured the magnetic-flux change rate in five flare events of different GOES classes using chromospheric/photospheric observations and compared its progression with observed nonthermal flare emission. We calculated the cumulated positive and negative magnetic flux participating in the reconnection process, as well as the total reconnection flux. Finally, we investigated the relations between the total reconnection flux, the GOES class of the events, and the linear velocity of the flare-associated CMEs.} {Using high-cadence H$\\alpha$ and TRACE~1600~{\\AA} image time-series data and MDI/SOHO magnetograms, we measured the required observables (newly brightened flare area and magnetic-field strength inside this area). RHESSI and INTEGRAL hard X-ray time profiles in nonthermal energy bands were used as observable proxies for the flare-energy release rate.} {We detected strong temporal correlations between the derived magnetic-flux change rate and the observed nonthermal emission of all events. The cumulated positive and negative fluxes, with flux ratios of between 0.64 and 1.35, were almost equivalent to each other. Total reconnection fluxes ranged between $1.8 \\times 10^{21}$~Mx for the weakest event (GOES class B9.5) and $15.5\\times 10^{21}$~Mx for the most energetic one (GOES class X17.2). The amount of magnetic flux participating in the reconnection process was higher in more energetic events than in weaker ones. Flares with more reconnection flux were associated with faster CMEs.} {} ",
        "introduction": "Solar flares are intriguing, intensely-studied phenomena. Much progress has been made in understanding the processes that occur on the Sun when a flare occurs. However, these events remain a partly-unsolved mystery. It is generally accepted that the energy fueling a solar flare is magnetic energy, which is accumulated slowly and stored inside the corona. After its release by magnetic field reconnection, it is converted into the kinetic energy of fast particles, plasma flows, heat, and MHD shock waves. Unfortunately, our observations remain insensitive to the primary energy-release process, since the energy release site in the reconnecting current sheet cannot be resolved with our current technical equipment. To learn about the properties of the energy release site and the processes occurring there, we therefore depend on exploring the plasma responses to the original energy-release process.  One of the observable consequences of the energy release in the corona are the bright, separating flare ribbons, which can be seen in UV bands and chromospheric spectral lines, and are prominent in H$\\alpha$. They are located on each side of the magnetic polarity inversion line and move away from it during the course of the flare. This apparent motion can be explained by the most widely accepted model of erupting flares, the CSHKP model \\citep{car64,sturr66,hir74,kopp76}. According to this, the separating flare ribbons are the chromospheric signatures of the energy release in the corona. The energy, which has been released at the primary energy-release site, is transported downward into the chromosphere along the newly reconnected field lines, by fast particles and thermal conduction. When it is deposited at the two footpoints of the field lines, which are rooted inside opposite magnetic polarities, the chromosphere brightens, i.e., the so-called flare kernels are created, and many of these adjacent kernels form elongated flare ribbons. The reconnection site is located below the flare-associated erupting filament/CME, and, since the erupting structure rises continuously to higher coronal heights, the reconnection also occurs at successively higher altitudes, as the flare progresses. As a consequence, field lines with footpoints rooted further and further away from each other are involved in the reconnection process. Thus, the newly brightened areas of both ribbons are located further and further apart and we observe separating flare ribbons.  Since the coronal reconnection site is coupled with the chromosphere by the reconnected field lines, it is obvious that the coronal magnetic reconnection rate, i.e., the rate at which magnetic field lines are reconnected, is associated with the separating flare ribbons. Based on the CSHKP model, \\citet{forbes84} and \\citet{forbes00} considered the rate of photospheric magnetic flux-change $\\dot{\\varphi}$. Using magnetic flux conservation, they derived a quantitative estimate for the coronal reconnection rate or magnetic flux change rate $\\dot{\\varphi}$, respectively, that can be derived from observable quantities  \\begin{equation} \\dot{\\varphi}=\\frac{\\partial}{\\partial t} \\int B_n\\,da, \\end{equation} where $B_n$ is the photospheric magnetic-field strength component perpendicular to the solar surface inside the newly brightened area $da$, which is swept up by the flare ribbons. We denote by $\\dot{\\varphi}$ the rate at which magnetic flux from formerly separated magnetic domains is unified. The original CSHKP model is a 2D model, since it describes the evolution of a flare in a vertical plane. In the third dimension, namely, the direction of the polarity inversion line, translational symmetry is assumed. This assumption converts the model to a 2$\\frac{1}{2}$D model. It is likely, however, that the extension in the third dimension is not continuous, but rather highly fragmented into temporary magnetic islands. However, we note that calculating the magnetic flux change rate $\\dot{\\varphi}$ does not require the assumption of translational symmetry. Therefore, $\\dot{\\varphi}$ is valid in three dimensions.  To test the CSHKP model, an observable quantity is required that can be regarded as a proxy for the energy release rate or reconnection rate, respectively, which we are incapable of measuring directly. The fast electrons that spiral downward along the newly reconnected field lines, generate microwave gyrosynchrotron emission, and when they deposit their energy at the footpoints of the reconnected field lines, hard X-ray (HXR) emission is produced by the nonthermal bremsstrahlung of electrons being scattered off ions. Therefore, the observed microwave and HXR fluxes act as indicators of the rate of accelerated electrons, which carry a large fraction of the total energy released during a flare \\citep[e.g.,][]{hud91,dennis03}. Thus, microwave and HXR emission can be used as proxies for the flare energy-release rate. If the CSHKP model is applicable, the derived magnetic-flux change rate $\\dot{\\varphi}$ should exhibit a similar temporal evolution as the observed HXR and microwave flux, i.e., the maxima in the $\\dot{\\varphi}$-profile should coincide approximately co-temporally with peaks in the observed nonthermal flare-emission. Strong temporal correlations of observed nonthermal emission with derived quantities characterizing the magnetic reconnection process, such as the magnetic flux change rate, the electric field at the reconnecting X-point, or the Poynting flux, which is transported into the reconnection region, were found in several cases \\citep[e.g.,][]{asa04,qiu04,jing05,isobe05,lee06,temmer07,mikl07}.  It is known from observations that the energy release in a dynamical flare is also related to the kinematics of the associated CME \\citep[e.g.,][]{maricic07,temmer08}. Statistical studies indicate that CME parameters, such as the velocity and kinetic energy, are correlated with the soft X-ray (SXR) peak flux and the time-integrated SXR flux, i.e., the flare fluence \\citep[e.g.,][]{moon02,bojan05c}. \\citet{qiu05} reported that a greater amount of total reconnection flux is related to higher CME velocities.  In this paper, we present a detailed analysis of five well-observed erupting flare events of different GOES classes. We calculated $\\dot{\\varphi}$ to test the model prediction, namely, the co-temporal evolution of the derived magnetic-flux change rate and observed nonthermal flare emission. In addition, we determined the cumulated reconnection-flux profiles for the positive and negative magnetic polarity domain as a function of time, the ratio of these fluxes, as well as the total flux that had to have been reconnected until the end of the flare. Finally, we investigated the relations between the total reconnection flux, the GOES class of the events, and the linear velocity of the flare-associated CMEs.  The paper is structured as follows. Section~\\ref{sec:dobs} contains descriptions of the used data sets, Sect.~\\ref{sec:analysis} describes the methods that we have applied, Sect.~\\ref{sec:results} lists the results for each flare, and in Sect.~\\ref{sec:discussion} the results are summarized and discussed. ",
        "conclusions": "\\label{sec:discussion}   In the following, we list potential errors in the determination of the magnetic reconnection flux that may arise from the applied methods, as well as the technical equipment and data acquisition procedure. Afterwards, we discuss the physical implications of our results.  \\subsection{Potential errors}  Potentially significant measurement errors, which can result from transient bright non-ribbon features, such as cosmic rays, were accounted for by eliminating short-lived features from the images (see Sect.~\\ref{sec:dobs} for more detail). Disregarding transient features can result in an overestimation by 100\\% of the total flare area and 20\\% of the total reconnection flux, if an image time series is strongly corrupted by particle hits. Bright coronal loops may also overestimate the flare area, and thus, the reconnected flux, when counted among flare pixels. Therefore, we excluded these regions from the analysis by setting the corresponding pixels during the time range, when these brightenings occurred, to zero.  A crucial factor is the threshold value, which is used to discern flare pixels and non-flare pixels. \\citet{longcope07} reported a change of 25\\% in the measured flux, when experimenting with cutoff values. We also tested several cutoff values and found a change of 5~--~15\\% in the measured reconnection flux.  A further effect, which may cause an overestimation of the flare area, and thus, the reconnection flux, is caused by a limit to the amount of charge that each pixel of a CCD chip can store. If there is too much charge at a particular location on the CCD chip, i.e., if a pixel is saturated, it will overflow to its neighboring pixels, preferentially the pixels above and below the saturated pixel. The signal in those pixels is questionable, since they will usually contain spilled charge. In three of the events (2005 January 15, X2.6, 2006 July 06, M2.5, and 2007 May 19, B9.5), H$\\alpha$ images were saturated during the impulsive phase, and the area and reconnection flux may thus be overestimated.  All sources of uncertainty mentioned so far are related to the flare area. The second quantity required to determine magnetic reconnection rates and fluxes, namely, the magnetic field, is measured routinely only in the photosphere. Thus, photospheric line-of-sight magnetograms are commonly used for this purpose. However, since the magnetic field strength decreases from the photosphere to the chromosphere/transition region, where the flare area is measured, reconnection fluxes are overestimated, when the magnetic field is derived from photospheric magnetograms, as was done for the present paper. \\citet{qiu07} reported an overestimation of about 20\\% in the reconnection flux derived from photospheric magnetic fields compared to the flux derived using potential field extrapolation to a height of 2~Mm.  As for the magnetic field, a further potential source of error must be accounted for. \\citet{berger03} reported that MDI measurements underestimate fields stronger than 1700~G, and \\citet{qiu07} estimated an upper limit to possibly underestimated reconnection flux due to MDI saturation effects, assuming that saturation occurs at field strengths of 1700, 1600, and 1500~G. In two of the events analyzed in this paper (2003 October 28, X17.2 and 2005 January 15, X2.6), small parts of the flare ribbons entered regions with field strengths stronger than 1500~G. According to \\citet{qiu07}, this may result in an underestimate of a few percent, at the very maximum, of the total reconnection flux. To account for this effect, we followed the cross-calibration study of \\citet{berger03} and multiplied magnetic field values by a correction factor of 1.56.  The misalignment between UV/H$\\alpha$ images and MDI magnetograms is another source of uncertainty, which can be estimated by artificially offsetting the two sets of images. We did not experiment with intentionally misaligned images, but \\citet{longcope07} and \\citet{qiu05} reported that artificial misalignment by up to 4\\arcsec~can account for up to 10~--~20\\% of change in the measured flux. However, the co-alignment accuracy that we achieved was higher than 2\\arcsec.  Bearing in mind all the sources of uncertainty discussed above, it is difficult to provide an overall error estimate for the total reconnection flux in each of the five events analyzed in this paper. Effects of over- and underestimation may cancel each other to a certain degree. To be on the safe side, an overestimate of 30~--~40\\% should be taken into account, especially for events in which the fluxes were derived from saturated H$\\alpha$ images.  \\subsection{Consideration of the results in view of the standard flare/CME model and numerical MHD models}\\label{sec:discussion2}  We derived the magnetic-flux change rate from chromospheric/photospheric observations of five flare events (GOES classes B, M, and X) and compared it to the observed nonthermal flare emission. In addition, we calculated the cumulated positive and negative magnetic flux participating in the reconnection process, as well as the total reconnection flux $\\varphi_{\\rm{tot}}$. We also determined the correlation between the $\\varphi_{\\rm{tot}}$, GOES peak, and linear velocity of the flare-associated CMEs.  We found good temporal correlations between the derived magnetic-flux change-rate and observed nonthermal emission in all events, i.e., hard X-ray peaks were clearly reflected in the magnetic flux change-rate profiles, as expected from the standard model, although the relative height of the peaks was not always reproduced in the $\\dot{\\varphi}$-profiles. In two events, for which non-saturated TRACE image time series were available with particularly high cadence ($\\leq 23$~s), the HXR peaks appeared delayed by roughly 1~min. Considering the amount of this delay, in \\citet{mikl07}, we speculated that it might be related to the travel time of a reconnected field line from the diffusion region to the lower edge of the current sheet. However, since we found a delay of HXR peaks in only two events so far, no conclusions can be drawn at present. It is, however, interesting to note that \\citet{warmuth09} calculated a travel time of $92$~s in the 2003 November 18 flare, when comparing their model of shock drift acceleration at the reconnection outflow termination shock with observations.  The derived total reconnection fluxes ranged between $1.8 \\times 10^{21}$~Mx for the weakest event and $15.5\\times 10^{21}$~Mx for the most energetic one (cf.~Table~\\ref{tab:3}). According to the standard model, equal amounts of positive and negative magnetic flux should participate in the reconnection process. The cumulated positive and negative fluxes that we found were roughly balanced, with flux ratios ranging between 0.64 and 1.35. According to \\citet{qiu05}, flux ratios of between 0.5 and 2 can be regarded as a good flux balance, given the numerous uncertainties involved in the measurements.  Beside determining reconnection rates and fluxes, we also calculated the correlation between GOES peak flux and: (1) total flare area; (2) mean magnetic field strength inside this area; and (3) total reconnection flux. Although neither the flare area nor the magnetic field were correlated with the GOES class, the total reconnection flux was. However, we note that our sample size is small. Therefore, we cannot exclude the possibility that the other two parameters are also related to the GOES peak flux, and statistical studies have indeed found such correlations. \\citet{nagashima06} analyzed 77 C, M, and X flares and detected a threshold value in the magnetic field that increases with the GOES peak flux. For example, X-class flares occurred only when the mean magnetic field in the active region was stronger than $\\sim 100$~G. \\citet{su07a} analyzed a sample of 31 flares and found that for events of stronger average magnetic field strength, the GOES peak flux tended to be higher. In a subsample of 18 out of the 31 events, they also detected a positive correlation between the GOES class and both the flare area and the magnetic flux, the magnetic flux showing a much stronger correlation than the area or the magnetic field. \\citet{su07a} attributed this to the magnetic flux being the product of the other two parameters. Although the applied methods for determining the flare area differ from our approach~--~\\citet{nagashima06} used intensity thresholds in soft X-ray images, and \\citet{su07a} took a contour level in photospheric magnetograms of the flaring region as a basis for the calculation of $A$~--~our findings are in line with the explanation given by \\citet{su07a}, because they also indicate that in this context the combination of area and magnetic field inside this area is important. It is conceivable that we did not find the weaker correlations between GOES class and $A$ and $\\langle B\\,\\rangle$ due to the small size of our sample, but we did find the stronger correlation between GOES peak flux and $\\varphi_{\\rm{tot}}$.  In addition, we incorporated our flux results into a larger set of results, taken from \\citet{qiu05}, \\citet{qiu07}, and \\citet{longcope07}, and again found a significant correlation between GOES peak flux and total reconnection flux ($R=0.6$, confidence level greater than 99\\%). However, the correlation became smaller ($R=0.3$), and was significant only with an 80\\%-confidence level, when the most energetic events (GOES classes $\\geq \\rm{X}10$) were excluded. This suggests that the rare, most energetic flares strongly contribute to the correlation. Still, the results indicate that the amount of magnetic flux participating in the reconnection process is larger in more energetic events than in weaker ones, as theoretically expected. The more magnetic flux is reconnected in a flare, the more energy is released into fast particles and can subsequently be deposited into the chromosphere. This energy heats the chromospheric plasma, which then evaporates, fills the flare loops, and causes them to emit soft X-ray radiation, measured by the GOES satellites. The event is assigned to a particular class in the GOES classification scheme according to its maximum soft X-ray emission, which can be understood as a measure of the cumulated energy in the hot thermal plasma \\citep[e.g.,][]{veronig02,veronig05}.  We also found that flares with more reconnection flux are associated with faster CMEs. This correlation, which was also reported by \\citet{qiu05}, can be explained by a feedback relationship between the CME kinematics and the reconnection process in the flare \\citep{bojan05c,maricic07,temmer08}. According to the standard flare/CME model, 50\\% of the magnetic flux reconnected during a dynamical (two-ribbon) flare is transported upward to higher coronal heights and even interplanetary space. The reconnection process in the flare, which is initiated in the wake of the CME, affects the CME by: (1) reducing the net tension of the overlying magnetic field; (2) increasing the magnetic pressure below the erupting flux-rope; and (3) supplying the flux-rope with extra poloidal flux \\citep{maricic07}. This enhances and prolongs the acceleration of the flux-rope and therewith additionally drives the CME. Therefore, CMEs that are associated with dynamical flares are generally faster than those linked with eruptive filaments, and CMEs that are associated with more powerful flares, i.e., flares of higher GOES class or fluence (time-integrated SXR flux), are of higher median speed and kinetic energy than weak-flare-associated CMEs \\citep{moon02,bojan05c}. In other words, the higher the amount of flux reconnected in a flare, the higher the amount of flux going upward, resulting ultimately in faster CMEs.  Numerical MHD models dealing with the formation of a flux rope, its sudden eruption and escape from the Sun provide additional evidence of a feedback relationship between an erupting structure and magnetic reconnection. \\citet{cheng03} accomplished 2.5-dimensional resistive MHD simulations of the evolution of an initially closed arcade field configuration. By imposing a shear-increasing footpoint motion, magnetic reconnection occurs, creating a new flux rope. The toroidal flux originally contained in the arcade is redistributed into both the slowly rising flux rope and the underlying arcade field. Thus, the magnetic shear in the arcade is reduced after the flux rope formation, but increases again, if additional footpoint motion occurs. As soon as the magnetic shear exceeds a critical value, another reconnection event occurs and a new flux rope is formed. When magnetic reconnection continues, this newborn flux rope rises with increasing velocity, and finally merges with the previously created flux rope. This process of flux rope formation and merging can be repeated as long as magnetic shear is continuously supplied. \\citet{cheng03} quantitatively compared the model results with low-corona observations of a filament eruption during a flare/CME event and found good agreement in the evolution of the flux rope height, velocity, and acceleration during the flux rope acceleration phase. Although the model does not address the escape of a flux rope from the Sun into interplanetary space, it shows that reconnection enhances the flux rope acceleration, and thus supports the scenario of a feedback relationship between the erupting structure and the reconnection in the flare.  \\citet{lin00} investigated the interaction between an already existing flux rope and magnetic reconnection. In their two-dimensional model, the flux rope experiences a catastrophic loss of equilibrium when the photospheric sources of the coronal magnetic field are brought together quasi-statically and reach a critical distance. Then, the flux rope jumps to a new equilibrium position at a higher altitude, and a vertical current sheet is created below the flux rope. The model shows that without reconnection in the current sheet the magnetic tension force is always strong enough to prevent the flux rope from reaching interplanetary space. However, even a modest reconnection rate (prescribed by the inflow Alfv\\'en Mach number $M_{\\rm{A}}$) is sufficient to allow its escape. Based on this loss-of-equilibrium model and its extension, which includes gravity in the calculations \\citep{lin04,reeves05b}, \\citet{reeves06} examined the relevance of reconnection for the flux-rope acceleration by analyzing its impact on the individual forces acting on the flux rope. They found that only the force due to the current sheet was affected considerably by changes in the reconnection rate $M_{\\rm{A}}$, while other forces, e.g., gravity, exhibited only minor changes with varying $M_{\\rm{A}}$. Slow reconnection rates resulted in longer current sheets, and thus, an increase in the downward force exerted by the current sheet on the flux rope. As a consequence, the total acceleration of the flux rope was decreased in cases of slow reconnection rates. When the reconnection rate was fast, the force due to the current sheet was diminished, because fast reconnection rates dissipate the current sheet. Therefore, the total flux rope acceleration was higher in cases of fast reconnection.  All of the models aforementioned are 2D or 2.5D models, which might be appropriate only in describing the eruption of a very long flux rope, especially when its height and width are much smaller than its length. In 3D models, the flux rope is anchored in the photosphere, and because of its final size, the overlying field can be pushed aside by the erupting structure. Thus, unlike in the 2D-approach, in 3D-model reconnection is not necessarily the main mechanism reducing the tension of the overlying field. However, \\citet{bojan08} demonstrated that in a 3D-flux-rope model without reconnection high accelerations are prevented because of the inductive decay of the current in the flux rope as the eruption progresses. On the other hand, including reconnection underneath the rising flux rope provides `fresh' poloidal flux to the rope. This preserves the current in the flux rope, and thus prolongs and enhances its acceleration.  The theoretical work discussed above reveals the importance of magnetic reconnection for the kinematics of the erupting structure. Reconnection in the current sheet occurs during the eruption of the flux-rope/CME, additionally drives the eruption, and may even be essential for a flux-rope to escape from the Sun, i.e., there may be no fast CMEs without magnetic reconnection, which is generally believed to play a major role in solar flares. However, it is known from observations that CMEs are not necessarily associated with flares. Instead, they may occur together with erupting filaments. After the eruption of a quiescent filament a growing system of post-eruption loops is often observed in the EUV range \\citep[e.g.,][]{bojan05c}. Morphologically, these loop systems are similar to postflare loops, except for the fact that they are not hot enough to be seen in soft X-rays. Quiescent filament eruptions occur in quiet regions, where the plasma-to-magnetic-pressure ratio $\\beta$ is generally larger than in active regions \\citep{wu01}. \\citet{bojan05} demonstrated that in a $\\beta > 0.1$ environment, the reconnection outflow region is not heated much. Therefore, an eruption taking place in such an environment should not be accompanied by a flare.  The loss-of-equilibrium model also indicates that insufficient plasma heating may be responsible for non-flare CMEs. \\citet{reeves05} simulated several erupting flux-rope/CME events, using background magnetic fields that ranged from weak to higher fields strengths. They found that weaker/higher background fields resulted in slower/faster CMEs, if all other parameters in the calculations remained the same. They also calculated GOES X-ray light curves for the 1~--~8~{\\AA} channel adding two different levels of soft X-ray background flux, namely, the average of the maximum/minimum 1~--~8~{\\AA} GOES flux during solar cycle~22. It turned out that with the weakest of the chosen background magnetic fields, adding the low X-ray background level produced an A5 flare, while adding the high X-ray background-flux, the GOES light curve hardly emerged from the background, i.e., in this case, the thermal energy release was not large enough to heat and evaporate sufficient plasma into the loops. Thus, the corresponding GOES light curve progression should probably not be considered to be a flare, although, magnetic reconnection was involved in the eruption of the flux rope/CME.  It is also known from observations that fast CMEs are not always associated with powerful flares \\citep[e.g,][]{bojan05c}. Although there are correlations between the kinetic energy/velocity of CMEs and the SXR peak flux in the corresponding flares \\citep[e.g.,][]{moon03,burkepile04}, CMEs of comparable speed can by all means be associated with flares of different GOES classes. Two of the flare/CME events analyzed in this paper provide an observational example. The CME-speeds on 2006 July 06 and 2007 May 19 were of the order of $900\\,\\rm{km\\,s^{-1}}$, but the GOES importance of the associated flares was M2.5 and~B9.5, respectively (cf.~Table~\\ref{tab:1}). The simulations carried out by \\citet{reeves05} also addressed this issue, and the results indicated that the CME speed is not indicative of the amount of associated X-ray flux, i.e., CMEs with very similar trajectories can have quite different flare responses. The authors simulated a high-mass flux rope erupting from a strong background field region as well as a low-mass flux-rope being ejected out of a weaker background field. The resulting velocity profiles for both events were similar. The flux-rope/CME speeds reached values of around $1100\\,\\rm{km\\,s^{-1}}$, although the corresponding GOES peak values, and thus, the importance of the associated simulated flares were different (C5 vs.~M2). \\citet{reeves05} attributed the difference in GOES class to the amount of thermal energy released in the current sheet in the model being directly related to the strength of the background magnetic field, while the kinetic energy of the flux rope depends on both the mass of the flux rope and the strength of the background field. Therefore, low mass CMEs from weak-field regions can have a similar velocity profile as massive CMEs from stronger field regions. The difference in the physical properties of both situations is evident only in the simulated X-ray emission of the associated flare.  Taking into account the results provided by numerical MHD simulations of solar eruptions, it becomes apparent that the combination of theoretical work and flare/CME observations is essential for improving our understanding of the underlying physical mechanisms and the relation between the two phenomena."
    },
    "0910/0910.0760.txt": {
        "abstract": "We report the detection of a dark substructure -- undetected in the HST-ACS F814W image -- in the gravitational lens galaxy SDSSJ0946+1006 (the ``Double Einstein Ring''), through direct gravitational imaging. The lens galaxy is of particular interest because of its relative high inferred fraction of dark matter inside the effective radius. The detection is based on a Bayesian grid reconstruction of the two-dimensional surface density of the galaxy inside an annulus around its Einstein radius. The detection of a small mass concentration in the surface density maps has a strong statistical significance. We confirm this detection by modeling the substructure with a tidally truncated pseudo-Jaffe density profile; in that case the substructure mass is $M_{\\rm{sub}}=(3.51\\pm 0.15)\\times 10^9\\msun$, located at $(-0.651\\pm0.038,1.040\\pm0.034)$'', precisely where also the surface density map shows a strong convergence peak (Bayes factor $\\Delta\\log{\\cal{E}}= -128.0$; equivalent to a $\\sim$16--$\\sigma$ detection). We set a lower limit of ${\\rm (M/L)}_{{\\rm V},\\odot}\\ga 120~ \\msun/{\\rm L}_{{\\rm V},\\odot}$ (3--$\\sigma$) inside a sphere of 0.3 kpc  centred on the substructure ($r_{\\rm tidal}$=1.1\\,kpc). The result is robust under substantial changes in the model and the data-set (e.g.\\ PSF, pixel number and scale, source and potential regularization, rotations and galaxy subtraction). It can therefore not be attributed to obvious systematic effects. Our detection implies a dark matter mass fraction at the radius of the inner Einstein ring of $f_{\\rm CDM}=2.15^{+2.05}_{-1.25}$ percent (68\\% C.L) in the mass range $4\\times10^6\\msun$ to $4\\times10^9\\msun$ assuming $\\alpha=1.9\\pm0.1$ (with $dN/dm\\propto m^{-\\alpha}$). Assuming a flat prior on $\\alpha$, between 1.0 and 3.0, increases this to $f_{\\rm CDM}=2.56^{+3.26}_{-1.50}$ percent (68\\% C.L). The likelihood ratio is 0.51 between our best value ($f_{\\rm CDM} = 0.0215$) and that from simulations ($f_{\\rm sim} \\approx 0.003$). Hence the inferred mass fraction, admittedly based on a single lens system, is large but still consistent with predictions. We expect to further tighten the substructure mass function (both fraction and slope), using the large number of systems found by SLACS and other surveys. ",
        "introduction": "In the process of building a coherent picture of galaxy formation and evolution, early-type galaxies play a crucial role. Often unfairly referred to as \\emph{red and dead} objects, many aspects about their structure and formation are still unknown. What is the  origin of the tight empirical relations between their global properties \\citep{Djorgovski87,Dressler87,Magorrian98,Ferrarese00,Gebhardt00,Bower92,Guzman92,Bender93}?  How do massive early-type assemble? What is the fraction of mass substructure populating the haloes of early-type galaxies and is this in agreement with the CDM paradigm \\citep{Kauffmann93,Moore99,Klypin99,Moore01,Maccio06,Diemand08,Springel08,Xu09}?  \\noi Gravitational lensing, especially in combination with other techniques, provides an invaluable and sometimes unique insight in answering these questions \\citep[e.g.][and references therein]{Rusin05,Treu04,Koopmans09}.  \\noi At the level of small mass structure lensing stands out as a unique investigative method; different aspects of the lensed images can be analysed to extract information about the  clumpy component of galactic haloes. Flux ratio anomalies, astrometric perturbations and time-delays, in multiple images of lensed quasars, can all be related to substructure at scales smaller than the images separation \\citep{Mao98,Bradac02, Chiba02, Dalal02, Metcalf02, Keeton03, Kochanek04, Bradac04, Keeton05,McKean07,Chen07,2009MNRAS.394..174M}.  \\noi As described by \\citet{Koopmans05} and \\citet{Vegetti09a}, the information contained in multiple images and Einstein rings of extended sources, can also be used. While the former three approaches only provide a statistical measure of the lens clumpiness, the latter allows one to identify and quantify of each single substructure, measuring for each of them the mass and the position on the lens plane. Both approaches are however complementary in that the former is more sensitive to low-mass perturbations, which are potentially present in large numbers, whereas the latter is sensitive to the rarer larger scale perturbations.  \\noi The method of direct gravitational imaging of the lens potential -- shortly described in the following section -- represents an objective approach to detect dark and luminous substructures in individual lens systems and allows on to statistically constrain the fraction of galactic satellites in early-type galaxies. Extensively described and tested in \\citet{Vegetti09a}, the procedure is here applied to the study of the double ring SDSSJ0946+1006 \\citep{Gavazzi08} from the sample of the \\emph{Sloan Lens ACS Survey }(SLACS), yielding the first detection of a dwarf satellite through its gravitational effect only, beyond the Local Universe.  \\noi The layout of the paper is as follows: in Section 2 we provide a general description of the method. In Section 3 we introduce the analysed data and in Section 4 we discuss the results of the modelling under the assumption of a smooth lens potential.  In Section 5 we describe the detection of the substructure. In Section 6 we present the error analysis and the model ranking. In Section 7 we discuss implication for the CDM paradigm and in Section 8 we conclude. Through out the paper we assume the following cosmological parameters with $H_0 = 73\\,\\kms\\Mpc^{-1}$, $\\Omega_{\\rm m}=0.25$ and $\\Omega_\\Lambda=0.75$. ",
        "conclusions": "We have applied our new Bayesian and adaptive-grid method for pixelized source {\\sl and} lens-potential modeling \\citep{Vegetti09a} to the analysis of HST data of the double Einstein ring system SLACS SDSSJ0946+1006 \\citep{Gavazzi08}. This system was chosen based on its large expected dark-matter mass fraction near the Einstein radius and the high signal-to-noise ratio of the lensed images. Although these two facts should be uncorrelated to the mass fraction of CDM substructure, both incidences maximize the change of detection.  We find that a smooth elliptical power-law model of the system leaves significant residuals near or above the noise level; these residuals are correlated and spread over a significant part of the lensed images. Through a careful modeling of this data including either lens-potential corrections or an additional (low-mass) simply parametrized mass component, we conclude that the massive early-type lens galaxy of SLACS SDSSJ0946+1006 hosts a large mass-to-light ratio substructure with a mass around $M_{\\rm{sub}}\\sim 3.5 \\times 10^9\\msun$, situated on one of the lensed images. A careful statistical analysis of the image residuals, as well as a number of more drastic robustness tests (e.g.\\ changing the PSF, pixel number and scale, regularization level and form, galaxy subtraction and image rotation), confirm and support this detection. Based on this detection, the first of its kind, we derive a projected CDM substructure mass fraction of $\\sim 2.2\\%$ for the inner regions of the galaxy, using the Bayesian method of \\citet{Vegetti09b}; this fraction is high, but still consistent with expectations from numerical simulations due to the large (Poisson) error based on a single detection.  The numerical details of our results can be further summarized as follows:  \\begin{enumerate}  \\item[{\\bf (1)}] Using a Bayesian {\\tt Multinest} Markov-Chain exploration of the full model parameter space, we show that the identified object has a mass of $M_{\\rm{sub}}=\\left(3.51\\pm0.15\\right)\\times 10^9\\msun$  (68\\% C.L.) and is located near the inner Einstein ring at $(-0.651\\pm0.038,1.040\\pm0.034)''$. The Bayesian evidence is in favor a model that includes a substructure versus a smooth elliptical power-law only, by $\\Delta \\log({\\cal E})=-128.2$. This is roughly equivalent to a $16-\\sigma$ detection.  \\medskip  \\item[{\\bf (2)}] At the redshift of the lens-galaxy a 3--$\\sigma$ upper limit in luminosity is found of $5.0\\times 10^{6}~{\\rm L}_{V,\\odot}$. Within the inner 0.3 kpc  and 0.6 kpc,  we find that the integrated mass is respectively $(5.8\\pm 0.3)\\times 10^8\\msun$ and $(1.1\\pm 0.05)\\times 10^9\\msun$ and hence lower limits of ${\\rm (M/L)}_{{\\rm V},\\odot}\\ga 120~ \\msun/{\\rm L}_{{\\rm V},\\odot}$ (3--$\\sigma$) and ${\\rm (M/L)}_{{\\rm V},\\odot}\\ga 218 ~\\msun/{\\rm L}_{{\\rm V},\\odot}$ (3--$\\sigma$). This is higher than of MW satellites, but maybe not unexpected for satellites near massive elliptical galaxies.  \\medskip  \\item[{\\bf (3)}]  The CDM mass fraction and a mass function slope are equal to  $f=2.15_{-1.25}^{+2.05}\\%$ and $\\alpha=1.88_{-0.10}^{+0.10}$, respectively,  at a 68\\% confidence level for a Gaussian prior on $\\alpha$ centred on $1.90\\pm0.1$. For a flat prior on $\\alpha$ between 1.0 and 3.0, we find $f=2.56_{-1.50}^{+3.26}\\%$. Asking whether the $f=2.15\\%$ is consistent with $f=0.3\\%$, we find a likelihood ratio of 0.51; indeed both are consistent. This is the result of the considerable measurement error found for $f$, because it is based on only a single detection.  \\end{enumerate}  \\noi This is the first application of our adaptive ``gravitational imaging'' method to real data and clearly shows its promise. In the near future we will apply the method to a larger set of SLACS lenses in order to constrain, via the statistical formalism presented in \\citet{Vegetti09b}, the dark matter fraction in substructure and the substructure mass function.  \\vfill"
    },
    "0910/0910.3642_arXiv.txt": {
        "abstract": "We identify and investigate the nature of the 20 brightest 250\\,${\\rm \\mu m}$ sources detected by the Balloon-borne Large Aperture Submillimetre Telescope (BLAST) within the central 150\\,arcmin$^2$ of the GOODS-South field. Aided by the available deep VLA 1.4\\,GHz radio imaging, reaching $S_{\\rm 1.4} \\simeq 40\\,{\\rm \\mu Jy}$ (4-$\\sigma$), we have identified radio counterparts for 17/20 of the 250\\,${\\rm \\mu m}$ sources. The resulting enhanced positional accuracy of $\\simeq 1$\\,arcsec has then allowed us to exploit the deep optical ({\\it HST}), near-infrared (VLT) and mid-infrared ({\\it Spitzer}) imaging of GOODS-South to establish secure galaxy counterparts for the 17 radio-identified sources, and plausible galaxy candidates for the 3 radio-unidentified sources. Confusion is a serious issue for this deep BLAST 250\\,${\\rm \\mu m}$ survey, due to the large size of the beam. Nevertheless, we argue that our chosen counterparts are significant, and often dominant contributors to the measured BLAST flux densities. For all of these 20 galaxies we have been able to determine spectroscopic (8) or photometric (12) redshifts. The result is the first near-complete redshift distribution for a deep 250\\,${\\rm \\mu m}$-selected galaxy sample. This reveals that 250\\,${\\rm \\mu m}$ surveys reaching detection limits of $\\simeq 40$\\,mJy have a median redshift $z \\simeq 1$, and contain not only low-redshift spirals/LIRGs, but also the extreme $z \\simeq 2$ dust-enshrouded starburst galaxies previously discovered at sub-millimetre wavelengths. Inspection of the LABOCA 870\\,${\\rm \\mu m}$ imaging of GOODS-South yields detections of $\\simeq 1/3$ of the proposed BLAST sources (all at $z > 1.5$), and reveals 250/870\\,${\\rm \\mu m}$ flux-density ratios consistent with a standard 40\\,K modified black-body fit with a dust emissivity index $\\beta = 1.5$. Based on their IRAC colours, we find that virtually all of the BLAST galaxy identifications appear better described as analogues of the M82 starburst galaxy, or Sc star-forming discs rather than highly obscured ULIRGs. This is perhaps as expected at low redshift, where the 250\\,$\\mu m$ BLAST selection function is biased towards spectral energy distributions which peak longward of $\\lambda_{rest} = 100\\, {\\rm \\mu m}$. However, it also appears largely true at $z \\simeq 2$. ",
        "introduction": "\\label{INTRO}  The observed far-infrared background peaks around a wavelength  $\\lambda \\simeq 200-250\\,{\\rm \\mu m}$ (Puget et al.~1996; Fixsen et al.~1998). Over the past 10 years, deep extragalactic surveys have closed in on this wavelength regime from both shorter and longer wavelengths, and there has been dramatic progress in uncovering populations of sources which undoubtedly {\\it contribute} to this background. Specifically, at wavelengths $\\simeq 4$ times longer than the peak ($\\lambda \\simeq 850 - 1200\\,{\\rm \\mu m}$), surveys with the Sub-millimetre Common-User Bolometer Array (SCUBA), Mambo, AzTEC and now LABOCA (e.g. Coppin et al., 2006; Bertoldi et al. 2007; Weiss et al. 2009; Austermann et al. 2010) have been used to successfully assemble substantial samples of submm-selected extragalactic sources, and the spectral energy distributions (SEDs) of these sources are certainly rising steeply towards shorter wavelengths. Conversely, at wavelengths $\\simeq 4$ times shorter than the peak, the {\\it Spitzer Space Telescope} has proved effective at extragalactic source selection up to wavelengths $\\lambda \\simeq 70\\,{\\rm \\mu m}$ (Magnelli et al., 2009). But until now, the effective production of deep extragalactic samples in the central wavelength range $\\lambda \\simeq 100 - 500\\,{\\rm \\mu m}$ has been prohibited by the atmosphere and/or limitations on instrument sensitivity and telescope aperture.  This situation should shortly be transformed by the SPIRE instrument on the 3.5-m diameter {\\it Herschel Space Observatory} (Griffin et al.~2007). However, a powerful first insight into the nature of the $250\\,{\\rm \\mu m}$-selected galaxy population is already being provided by the the Balloon-borne Large Aperture Submillimetre Telescope (BLAST).  BLAST is a 1.8\\,m diameter stratospheric balloon telescope that operates at an altitude of approximately 35\\,km, above most of the atmospheric water vapour which essentially prohibits effective far-infrared observations from the ground. BLAST is thus half the size of {\\it Herschel}, and it is equipped with a camera which is a prototype of the SPIRE camera, enabling simultaneous broad-band imaging with central wavelengths of 250, 350 and 500\\,${\\rm \\mu m}$. In 2006, BLAST undertook an 11-day flight from Antarctica, during which it executed an observing programme which included the first deep far-infrared imaging survey ever undertaken within the Extended {\\it Chandra} Deep Field South (ECDFS).  This unique survey (described in more detail in Section 2) has already been the subject of several investigations. First results on the far-infrared number counts and the resolution of the background were presented by Devlin et al. (2009), with more detailed investigations of these two key topics being presented by Patanchon et al. (2009) and Marsden et al. (2009). Pascale et al. (2009) used the survey to constrain cosmic star-formation history, while Viero et al. (2009) studied correlations in the background. Dye et al. (2009) and Ivison et al. (2010) used the supporting radio (VLA) and mid-infrared {\\it Spitzer} imaging of the field to attempt to identify countparts to the brighter BLAST sources, and to explore the far-infrared/radio correlation in distant galaxies.  The study by Dye et al. (2009) showed that the BLAST beam, at least at $250\\, {\\rm \\mu m}$, is still small enough to allow the successful identification of a reasonable fraction of BLAST sources with radio and mid-infrared counterparts. However, the supporting data over the full 9\\,deg$^2$ BLAST survey area are of variable quality, and Dye et al. attempted to secure identifications for sources selected at all 3 BLAST wavelengths. As a result, they succeeded in identifying counterparts for only $\\simeq 55$\\% of the sources in the BLAST catalogues detected at $> 5\\,\\sigma$ in at least one waveband. This inevitably limits the conclusions that can be drawn about the redshift distribution of the BLAST source population, although Dye et al. concluded that $\\simeq 75$\\% of the sources in the shallow+deep catalogue lie at $z < 1$.  The aim of the study presented here is to establish a near-complete redshift distribution for a subset of the BLAST $250\\, {\\rm \\mu m}$ sources selected down to a fainter flux-density limit of $S_{250} \\simeq 40$\\,mJy, and subsequently to establish the basic physical properties (stellar mass, size, morphology, star-formation history, AGN content) of the galaxies which host the observed $250\\, {\\rm \\mu m}$ emission. To try to achieve this, we have deliberately confined our attention to the small sub-area of the BLAST map which contains the very best optical ({\\it HST} ACS), near-infrared (VLT ISAAC) and mid-infrared ({\\it Spitzer}) imaging data, and the highest density of optical spectroscopic redshifts. This area, the $\\simeq 150$\\,arcmin$^2$ GOODS-South field (Dickinson et al. 2004), is also where the radio (VLA) imaging is most sensitive, and it has been mapped in its entirety at $870\\, {\\rm \\mu m}$ by LABOCA (Weiss et al. 2009), and at 1.1\\,mm by AzTEC (Scott et al. 2010). GOODS-South is also the chosen location for some of the first deep high-resolution near-infrared imaging currently being undertaken with WFC3 on the refurbished {\\it HST}, and is the future target of planned deep sub-mm and far-infrared imaging with SCUBA2 (Holland et al. 2006) and SPIRE+PACS on {\\it Herschel}.  Thus, by confining our attention to GOODS-South we can explore just what can be achieved given the best possible supporting data. Conversely, we can also establish what level of supporting data is actually {\\it required} for an effective, complete study of the future, larger, $250\\, {\\rm \\mu m}$-selected galaxy samples which will be produced by {\\it Herschel}. In addition, the wealth of existing information on other galaxy populations in this field (see Section 2) facilitates the comparison of BLAST galaxies with reference samples of field galaxies selected at similar redshifts and/or stellar masses, and to explore the frequency of mergers. Finally, the recent acquisition of deep sub-millimetre and millimetre wavelength imaging in the same field allows a first exploration of how the $250\\, {\\rm \\mu m}$ and $870/1100\\, {\\rm \\mu m}$ populations are related.  To assemble a useful sample of $250\\, {\\rm \\mu m}$ sources within this field, we have pushed the BLAST source selection significance threshold to $\\sim 3.5\\sigma$. However, the resulting sample is still sufficiently small (20 sources) that we can afford to explore the properties of individual sources in some detail. As a result this work should be regarded as a pilot study of the mix of source populations which can be uncovered in a deep $250\\, {\\rm \\mu m}$ survey, and an exploration of how best to overcome the problems that are encountered in trying to identify and study the faint far-infrared sources in maps which are inevitably highly confused (note that the {\\it Herschel} beam at $500\\, {\\rm \\mu m}$ is essentially the same size as the BLAST beam at $250\\, {\\rm \\mu m}$). The broader statistical robustness of the results presented here is clearly limited both by small number statistics, and by cosmic variance.  The layout of the paper is as follows. In Section 2 we describe the selection and robustness of the BLAST $250\\, {\\rm \\mu m}$ sub-sample within the central region of GOODS-South, and then summarize the wealth of available multi-frequency imaging and spectroscopy in the field. In Section 3 we describe the process of obtaining radio and optical counterparts for the BLAST sources, and highlight the care required to overcome the problems associated with source blending in the confused BLAST maps. In Section 4 we provide or derive redshifts for all of the sources, including secondary/alternative counterparts. Then, in Section 5, we present multi-colour images of the BLAST galaxies and provide notes on each of the individual sources. Finally, in Section 6 we review our results and discuss the robustness of our conclusions. In particular we present and discuss the implications of the striking associations found between the BLAST $250\\, {\\rm \\mu m}$ sources and the LABOCA  $870\\, {\\rm \\mu m}$ sources recently uncovered by Weiss et al. (2009). We also discuss the robustness of our derived redshift distribution for $250\\, {\\rm \\mu m}$-selected sources, and briefly compare the mid-infrared colours of BLAST galaxies with those of other known galaxies at both high and low redshifts. Our conclusions are summarized in Section 7. ",
        "conclusions": "\\label{DISCUSSION}  \\begin{table} \\begin{tabular}{rrccrl} \\hline BLAST & LABOCA & $d_{B-L}$ & $d_{L-Rad}$ & \\phantom{00}$P_{LB}$ & Notes\\\\ ID\\phantom{ST}& ID\\phantom{ST}  & /arcsec  & /arcsec & & \\\\ \\hline 59 &  10+34  &  8.9  & 7.7  &   0.004  &Blend\\\\ 66 &  18     &  7.5  & 5.5  &   0.008   \\\\ 158&  79     &  7.4  & 5.0  &   0.011   \\\\ 318&  67     &  12.4 & 3.7  &   0.019   \\\\ 593&  12     &  8.7  & 2.4  &   0.008\\\\ 732&  32     & 9.6   & 4.4  &   0.012 & Opt ID\\\\ \\hline \\end{tabular} \\caption{Associations between BLAST 250\\,${\\rm \\mu m}$ and LABOCA LESS 870\\,${\\rm \\mu m}$ sources within the central 150\\,arcmin$^2$ of GOODS-South. $d_{B-L}$ is the distance, in arcsec, between the BLAST and LABOCA source. For comparison, $d_{L-Rad}$ is the distance, in arcsec, between the LABOCA source and the radio identification listed in Table 2. $P_{LB}$ is the probability that the association between the BLAST and LABOCA source is the result of chance, calculated in an analogous manner to the $P$ values calculated for the BLAST-Radio associations in Table 2, but with a search radius of 18\\,arcsec.} \\end{table}  \\subsection{LABOCA sub-mm detections and 250/870\\,${\\rm \\mu m}$ flux ratios}  Within the central area of GOODS-South under study here, Weiss et al. (2009) have extracted 10 LABOCA 870\\,${\\rm \\mu m}$ sources with a raw significance $> 4\\sigma$. Searching around the BLAST positions out to a search radius of 18\\,arcsec, we find that seven of these 870\\,${\\rm \\mu m}$ sources coincide with 250\\,${\\rm \\mu m}$ sources to within a positional accuracy of $< 13$\\,arcsec (including two associated with BLAST 59). These associations are tabulated in Table 4. We adopted a search radius of 18\\,arcsec by adding the positional uncertainty adopted for the BLAST sources in Section 3.1.1 ($\\sigma_{pos} = 6$\\,arcsec) in quadrature with an assumed positional uncertainty in the LABOCA sources of $\\sigma_{pos} = 4$\\,arcsec (calculated assuming the 19\\,arcsec LABOCA beam, and a typical deboosted S/N = 3), and then multiplying by 2.5 to ensure 95\\% completeness.  The LABOCA and BLAST source surface densities are both so much lower than the radio, mid-infrared or optical source surface densities, that all of these associations are statistically compelling, as demonstrated by the derived values of $P_{LB}$ given in Table 4. The other 3 LABOCA sources in the field are not near to any BLAST source, and so the choice of search radius appears to have been sensible.  \\begin{figure} \\begin{tabular}{c} \\includegraphics[width=0.46\\textwidth]{870.ps}\\\\ \\\\ \\\\ \\includegraphics[width=0.46\\textwidth]{250_870.ps}\\\\ \\\\ \\\\ \\end{tabular} \\caption{The LABOCA LESS 870\\,${\\rm \\mu m}$ detections and non-detections of the BLAST 250\\,${\\rm \\mu m}$ sources. The upper panel shows 870\\,${\\rm \\mu m}$ flux density plotted against redshift (spectroscopic or photometric) as derived from the primary proposed galaxy identification for each BLAST source. All the 870\\,${\\rm \\mu m}$ detections are confined to the high-redshift BLAST subsample. The non-detections are indicated by arrows plotted at the adopted conservative $\\simeq 3$-$\\sigma$ flux-density limit of $S_{870} < 4\\, {\\rm mJy}$. The lower panel plots the derived 250/870\\,${\\rm \\mu m}$ flux-density ratios (or limits), again versus redshift. The anticipated trend of declining 250/870\\,${\\rm \\mu m}$ flux-density ratio with redshift is clearly evident. The solid line, derived by redshifting a simple modified black-body spectrum with $T=40$\\,K, and $\\beta = 1.5$, provides a very good description of the data. The dotted curve shows the effect of increasing temperature to $T = 45$\\,K, and the dashed line the effect of reducing it to $T = 35$\\,K. In calculating the observed flux ratios, neither the BLAST nor the LABOCA flux densities have been deboosted.} \\end{figure}  The associations listed in Table 4 provide an opportunity to check that our adopted positional uncertainties are indeed reasonable. First, the median value of $d_{L-Rad} = 5.0$ implies that (assuming zero uncertainty in the radio positions), the positional uncertainty in the LABOCA positions is $\\sigma_{pos} \\simeq 4.5$ arcsec (allowing for the radial weighting in the observed distribution), showing that our adoption of 4\\,arcsec was at least approximately correct. The median value of the BLAST-LABOCA positional offset is $d_{B-L} = 8.9$, implying $\\sigma_{B-L} \\simeq 7.5$\\,arcsec. Subtracting 4.5 in quadrature from 7.5, leads to the conclusion that $\\sigma_{pos} = 6$\\,arcsec for the BLAST sources, as adopted in Section 3.1.1. Repeating this analysis with means instead of medians, yields $\\sigma_{pos} = 6.5$\\,arcsec.  Working backwards, this empirical check on the uncertainty in the positions of the BLAST sources considered in this paper re-affirms that they are indeed genuine $\\simeq 4\\sigma$ sources prior to flux deboosting (including the contribution of confusion to the noise).  The more scientifically interesting aspect of these LABOCA detections is that all of them are associated with {\\it high-redshift} BLAST sources. The 870\\,${\\rm \\mu m}$ flux densities of the LABOCA detections are listed in Table 3. From examination of this Table it can be seen that the LABOCA detections are confined to proposed identifications with $z > 1.5$. This point is illustrated in the upper panel of Fig. 9, where we have also adopted a conservative $\\simeq 3\\sigma$ flux-density limit of $S_{870} < 4\\, {\\rm mJy}$ for the non-detections. Not surprisingly, it can be seen that the redshift distribution of the LABOCA-detected BLAST sources is consistent with that which is displayed by galaxy samples selected at 850$\\mu m$ with SCUBA.  The 870$\\mu m$ detections thus provide strong, independent support for most of our proposed high-redshift BLAST identifications, especially since the identifications and redshifts were established without any reference to the LABOCA data. However, given the existence of the LABOCA data, one can then also ask whether the LABOCA {\\it non}-detections of 4 of our proposed $z > 1.5$ identifications cast doubt on their reality. To check this we have therefore plotted, in the lower panel of Fig. 9, the 250/870\\,${\\rm \\mu m}$ flux-ratios of the BLAST galaxies versus redshift. Here, despite the fact that only lower limits are available at low redshift, the anticipated trend of declining 250/870\\,${\\rm \\mu m}$ flux-density ratio with redshift is evident. It is also clear that, with one exception (BLAST 654), the non detection of the high-redshift candidates is not sufficiently deep to cast serious doubt on the proposed redshifts.    The solid curve shown in the lower panel of Fig. 9 is the predicted decline in colour with increasing redshift for a simple modified black-body spectrum with an assumed rest-frame temperature $T = 40$\\,K and dust emissivity emissivity index $\\beta = 1.5$. This offers a remarkably good description of the high-redshift data and is at least consistent with the lower redshift limits. In other words, in retrospect we should not be surprised that the 250\\,${\\rm \\mu m}$-selected galaxies at $z < 1$ have not been detected in the LABOCA LESS map, nor that the deep 250\\,${\\rm \\mu m}$ imaging analysed here has reached sufficient depth to allow us to detect the sub-mm galaxy population at $z \\simeq 2$. The dashed and dotted curves simply illustrate the effect of increasing or decreasing the assumed dust temperature by 5\\,K, keeping $\\beta = 1.5$. Given the mismatch between the BLAST and LABOCA beams, and the fact that neither set of flux densities has been deboosted in constructing Fig. 9, we do not want to over-emphasize the limited extent to which one can determine the temperature of the dust emission from the 250/870\\,${\\rm \\mu m}$ flux-ratios of the BLAST galaxies. However, at the very least, Fig. 9 provides a basic sanity check on the plausibility of our inferred galaxy identifications and redshifts.  \\subsection{The redshift distribution for the 250\\,${\\rm \\mu m}$ sample}  \\begin{figure} \\begin{tabular}{c} \\includegraphics[width=0.46\\textwidth]{redshift_hist.ps}\\\\ \\\\ \\\\ \\includegraphics[width=0.46\\textwidth]{redshift_hist2.ps}\\\\ \\\\ \\\\ \\end{tabular} \\caption{An exploration of the robustness of the basic form of the redshift distribution for the 250\\,${\\rm \\mu m}$ sources in GOODS-South. Sources 654 and 552 have been rejected from both distributions for the reasons detailed in the text. The upper panel shows the redshift distribution derived from the primary identifications (plus 59-2) listed in Table 3. In the lower panel we have replaced the primary identifications by the secondary identifications in the 5 appropriate cases, and in addition have removed the other source which lacks either a radio or a LABOCA counterpart (BLAST 193). Despite these 6 changes, the two distributions are statistically indistinguishable. In the upper panel there are 10/19 galaxies at $z > 1$. In the lower plot there are 8/18 sources at $z > 1$. A median redshift $z \\simeq 1$ is clearly a fair description of the data presented here, given the uncertainties in identifications and redshifts.} \\end{figure}  The inferred redshift distribution of our 20-source BLAST 250\\,${\\rm \\mu m}$ sample is presented in Fig. 10. We show two alternative versions, in order to illustrate the extent to which its basic form is robust against changes in ambiguous identifications. BLAST 654 and 552 are excluded from both plots, the former because its 250/870 \\,${\\rm \\mu m}$ limit appears anomalously high for its proposed redshift, the latter because its non detection with LABOCA, combined with the low significance of its radio identification and a non detection at 24\\,${\\rm \\mu m}$ casts serious doubt on the robustness of this identification. In the upper panel we plot the histogram derived from the primary galaxy identifications/redshifts listed in Table 3 (we include both 59-1 and 59-2). In this case 10/19 sources lie at $z > 1$, and the distribution appears almost bi-modal. In the lower panel we have replaced the primary identifications with the secondary/alternative galaxy ID for all potentially ambiguous cases (i.e. BLAST sources 6, 257, 503, 732, 830) and have also simply ejected BLAST 193 from the sample (because it possesses neither a VLA radio identification, nor a LABOCA sub-millimetre detection). This more sceptical approach still leaves 8/18 sources at $z > 1$, of which 6 have been detected by LABOCA.    \\begin{figure*} \\includegraphics[width=1.0\\textwidth]{IRCOLOR.ps} \\caption{The location of the adopted primary (red circles) and secondary (green stars) galaxy counterparts for the BLAST 250\\,${\\rm \\mu m}$ sources on the {\\it Spitzer} IRAC colour-colour plane. The small grey points indicate the positions of field galaxies within GOODS-South, as derived from all galaxies in the GOODS-MUSIC catalogue down to limiting IRAC AB magnitudes of $m_{3.6} = 24.0,\\ m_{4.5} = 23.4,\\ m_{5.8} = 22.0,\\ m_{8.0} = 22.0$. Each BLAST identification is labelled with its redshift. The three tracks are derived by simulating the effect of redshifting, from  $z = 0$ to $z = 3.5$, the spectral energy distributions of two well-known local star-forming galaxies (the Virgo Sc galaxy VCC1972, and M82) and the more highly obscured ULIRG Apr220. The blue solid track is for VCC1972, the green dashed track is for M82, and the red dot-dash track  is for Arp220, with the SEDs for all three galaxies taken from the multi-wavelength fits derived by Devriendt et al. (1999). The locations of the BLAST primary identifications are in general very well described by the VCC1972 or M82 tracks, which intially move rapidly in observed $5.8-8.0\\,{\\rm \\mu m}$ colour as the 8${\\rm \\mu m}$ PAH feature moves out of the longest IRAC filter, and later move rapidly up in $3.6-4.5\\,{\\rm \\mu m}$ colour as the 1.6${\\rm \\mu m}$ peak in the starlight moves into the $4.5\\,{\\rm \\mu m}$ filter. Very few of the BLAST galaxies lie on the Arp220/ULIRG track, with the vast majority being far better described by the less reddened star-forming galaxies. At low redshift this tallies well with the appearance of the BLAST identifications shown in Fig. 7, which shows that they are generally isolated dusty star-forming discs. Moreover, certainly at $z < 1$, 250\\,${\\rm \\mu m}$ selection, especially over the relatively small field of GOODS-South will be biased towards cooler SEDs which peak longwards of $\\lambda = 100\\, {\\rm \\mu m}$. This is true for star-forming disc galaxies, but not true for the hotter ULIRGs selected at 60\\,${\\rm \\mu m}$ with {\\it IRAS}. Perhaps more surprising (given their `ULIRG-like' star-formation rates) is the fact that the $z \\simeq 2$ BLAST/LABOCA galaxies also appear relatively unobscured at IRAC wavelengths.}  \\end{figure*}  For our sample, these two alternative redshift distributions are statistically indistinguishable, and we conclude that, at least in GOODS-South, $\\simeq 50$\\% of the 250\\,${\\rm \\mu m}$ sources with $S_{250} > 35\\,{\\rm mJy}$ lie at $z > 1$. Any evidence for bimodality is statistically insignificant, but the median redshift, and the fact that $\\simeq 1/3$ of the sample lies at $z > 1.5$, appear to be robust results. The precise form of the distribution is of course aflicted here both by small-number statistics, and by the potential effects of cosmic variance. It will therefore be interesting  to see the extent to which the redshift distribution derived here applies to the much larger samples of 250\\,${\\rm \\mu m}$ sources which should be uncovered with {\\it Herschel}.  It is interesting to assess how our results on this small, deep, near-complete BLAST subsample compare with those reported for the full 5-$\\sigma$ BLAST sample of the wider 9 degree$^2$ of the Extended Chandra Deep Field South by Dye et al. (2009). This comparison must be done with care, and the limitations understood, because Dye et al. were studying a larger but brighter sample, with on average much poorer supporting data. Dye et al. attempted both radio and 24\\,${\\rm \\mu m}$ identifications for 350 BLAST sources, secured identifications for 198 of them, and in the end secured redshifts for 74 of these (i.e. for $\\simeq $1/5th of the sample). It is not clear what conclusions can be drawn from the resulting redshift distribution, given the low-level of completeness and the mix of flux-density limits.  More instructive for the present comparison is the redshift distribution derived by Dye et al. for the $>5\\sigma$ BLAST sources in the central region of the ECDFS, with the best supporting data from the FIDEL survey. Here Dye et al. identified 50\\% of the BLAST sources ($\\simeq 75$\\% of the 250\\,${\\rm \\mu m}$-selected sources) at either radio or mid-infrared wavelengths. The median redshift reported in this sub-area by Dye et al. is $z_{med} \\simeq 0.95$. Given that the redshift information includes some COMBO 17 redshifts which (as discussed here for, e.g. BLAST 66) are undoubtedly under-estimates, given that the Dye et al. redshift distribution is more incomplete than that presented here, and given that the effective 250\\,${\\rm \\mu m}$ flux density limit of the Dye et al. deep sample is still $>50$\\% higher than in the present study, we conclude there is no real evidence for any inconsistency between our redshift distribution and that derived by Dye et al. (2009). The only visible difference is that the high-redshift tail in the present deep study extends to somewhat higher redshifts than that reported by Dye et al., but this is unsurprising given the above-mentioned selection effects.   \\subsection{The mid-infrared colours of the 250\\,${\\rm \\mu m}$ galaxies.}  Finally, in Fig. 11 we briefly investigate the location of the galaxy counterparts for the BLAST 250\\,${\\rm \\mu m}$ sources on the {\\it Spitzer} IRAC colour-colour plane. This is of interest for a number of reasons. First, it is worthwhile to check how the mid-infrared colours of the BLAST-selected galaxies compare with those of the general galaxy population in GOODS-South. Second, we can compare their colours with expectations based on the redshifted SEDs of known local far-infrared bright galaxies. Third, it is of potential interest to check whether the counterparts of 250\\,${\\rm \\mu m}$-selected galaxies can, in future, be reliably isolated on the basis of IRAC colour (a point of importance given the potentially extensive sky coverage which could be provided by the {\\it Spitzer} warm mission).  We choose to plot $3.6\\,{\\rm \\mu m}-4.5\\,{\\rm \\mu m}$ colour versus $5.8\\,{\\rm \\mu m}-8.0\\,{\\rm \\mu m}$ colour, as advocated by Stern et al. (2005), rather than $3.6\\,{\\rm \\mu m}-5.8\\,{\\rm \\mu m}$ colour versus $4.5\\,{\\rm \\mu m}-8.0\\,{\\rm \\mu m}$ colour as advocated by some other authors (e.g. Lacy et al. 2004; Pope et al. 2006). We do this to avoid unhelpful correlations, and because only $3.6\\,{\\rm \\mu m}-4.5\\,{\\rm \\mu m}$ colours will be available from the {\\it Spitzer} warm mission.  In Fig. 11 we plot the location of the adopted primary and secondary galaxy counterparts for the BLAST 250\\,${\\rm \\mu m}$ sources, superimposed on the positions of field galaxies within GOODS-South, as derived from all galaxies in the GOODS-MUSIC catalogue down to limiting IRAC AB magnitudes of $m_{3.6} = 24.0,\\ m_{4.5} = 23.4,\\ m_{5.8} = 22.0,\\ m_{8.0} = 22.0$. In addition, we show tracks in the colour-colour diagram produced by redshifting (from $z = 0$ to $z=3.5$) the spectral energy distributions of two well-known local star-forming galaxies (the Virgo Sc galaxy VCC1972, and the starburst galaxy M82), along with the analogous predicted track for the highly-obscured ULIRG Apr\\,220.  The main points we can learn from this diagram can be summarised as follows. First, the primary BLAST galaxy identifications are very well described by the two star-forming galaxy SEDs, with the red data points essentially mapping out the redshift evolution of the tracks. The rapid horizontal movement at low redshift ($z = 0 - 0.5$) is caused by the polycyclic aromatic hydrocarbon (PAH) features moving through the longer two IRAC bands. The subsequent rapid vertical movement from $3.6\\,{\\rm \\mu m}-4.5\\,{\\rm \\mu m} < 0$ to  $3.6\\,{\\rm \\mu m}-4.5\\,{\\rm \\mu m} > 0$ is the result of the stellar photospheric peak at $\\lambda_{rest} \\simeq 1.6 {\\rm \\mu m}$ moving through the shorter two IRAC bands at $z \\simeq 1.5$.  Second, it can be seen that very few of our primary identifications appear to be analogues of Arp\\,220, although at the very highest redshifts all three templates become indistinguishable. Our low-redshift BLAST galaxies stay relatively blue in $3.6\\,{\\rm \\mu m}-4.5\\,{\\rm \\mu m}$ colour, indicating their star-light is not swamped by the mid-infrared power-law produced by very hot dust (AGN or starburst heated), as is seen in Arp\\,220 and other local ULIRGs. Even at $z \\simeq 2$ most of the BLAST/LABOCA galaxies have colours expected from starlight (albeit in some cases reddened to $3.6\\,{\\rm \\mu m}-4.5\\,{\\rm \\mu m} > 0.25$). Our highest-redshift candidates do lie in the region of the diagram expected for $z \\simeq 3$, but at this point essentially all tracks converge on the same location in the diagram.  It is interesting to speculate why the Arp\\,220 track is such a poor description of the locus occupied by the BLAST sources. At low redshift it is perhaps unsurprising, as starburst galaxies like M82 and Sc star-forming discs both have (relatively) cool SEDs which peak longward of $\\lambda_{rest} \\simeq 100\\,{\\rm \\mu m}$, whereas the ultra-luminous ULIRGS first uncovered by {\\it IRAS} at $60\\,{\\rm \\mu m}$ peak at shorter wavelengths. A $250\\,{\\rm \\mu m}$ survey is thus biased towards the selection of these cooler more `normal' starbursts at low redshift, especially since highly-obscured ULIRGs are rather rare objects. The appearance of the low-redshift BLAST galaxies in Fig. 7 (i.e. star-forming discs) is thus consistent with their IRAC colours.  Perhaps more surprising (given their `ULIRG-like' star-formation rates) is the fact that the $z \\simeq 2$ BLAST/LABOCA galaxies also appear relatively unobscured at IRAC wavelengths. However, this is arguably consistent with their 250/870\\,${\\rm \\mu m}$ flux-ratios as discussed in Section 6.1, which appear to be consistent with a relatively cool dust temperature of 40\\,K (see Fig. 9).  Finally we revisit the claim made by Yun et al. (2008) that a simple and robust sub-mm galaxy candidate selection criterion is provided by $3.6\\,{\\rm \\mu m} - 4.5\\,{\\rm \\mu m} > -0.2$. It is obvious from Fig. 11 that this criterion is in fact satisfied by essentially any galaxy at $z > 1.5$, and thus cannot be used to isolate any special sub-class of source at these redshifts.  These issues, along with the other physical characteristics of the BLAST galaxies (size, morphology, stellar mass, star-formation history) will be explored in more detail by Targett et al. (in preparation)."
    },
    "0910/0910.3197_arXiv.txt": {
        "abstract": "We study the statistical distributions of the spins of generic black-hole binaries during the inspiral and merger, as well as the distributions of the remnant mass, spin, and recoil velocity.  For the inspiral regime, we start with a random uniform distribution of spin directions $\\vec{S}_1$ and $\\vec{S}_2$ over the sphere and magnitudes $|\\vec{S}_1/m_1^2|=|\\vec{S}_2/m_2^2|=0.97$ for different mass ratios, where $\\vec S_i$ and $m_i$ are the spin-angular momentum and mass of the $i$th black hole. Starting from a fiducial initial separation of $r_i=50M$, we perform 3.5-post-Newtonian-order evolutions down to a separation of $r_f=5M$, where $M=m_1+m_2$, the total mass of the system. At this final fiducial separation, we compute the angular distribution of the spins with respect to the final orbital angular momentum, $\\vec L$. We perform $16^4=65536$ simulations for six mass ratios between $q=1$ and $q=1/16$ and compute the distribution of the angles $\\hat{\\vec{L}}\\cdot\\hat{\\vec{\\Delta}}$ and $\\hat{\\vec{L}}\\cdot\\hat{\\vec{S}}$, directly related to recoil velocities and total angular momentum. We find a small but statistically significant bias of the distribution towards counter-alignment of both scalar products.  A post-Newtonian analysis shows that radiation-reaction-driven dissipative effects on the orbital angular momentum lead to this bias. To study the merger of black-hole binaries, we turn to full numerical techniques. In order to make use of the numerous simulations now available in the literature, we introduce empirical formulae to describe the final remnant black hole mass, spin, and recoil velocity for merging black-hole binaries with arbitrary mass ratios and spins. Our formulae are based on the post-Newtonian scaling, to model the plunge phase, with amplitude parameters chosen by a least-squares fit of recently available fully nonlinear numerical simulations, supplemented by inspiral losses from infinity to the ISCO. We then evaluate those formulae for randomly chosen directions of the individual spins and magnitudes as well as the binary's mass ratio. The number of evaluations has been chosen such that there are 10 configurations per each dimension of this parameter space, i.e. $10^7$. We found that the magnitude of the recoil velocity distribution decays exponentially as $P(v) \\sim \\exp(-v/2500\\ \\KMS)$ with mean velocity $<v> = 630\\ \\kms$ and standard deviation $\\sqrt{<v^2> - <v>^2} = 534\\ \\kms$, leading to a $23\\%$ probability of recoils larger than $1000\\ \\KMS$, and a highly peaked angular distribution along the final orbital axis. The studies of the distribution of the final black-hole spin magnitude show a universal distribution highly peaked at $S_f/m_f^2=0.73$ and a $25^\\circ$ misalignment with respect to the final orbital angular momentum, just prior to full merger of the holes. We also compute the statistical dependence of the magnitude of the recoil velocity with respect to the ejection angle. The spin and recoil velocity distributions are also displayed as a function of the mass ratio. We finally also compute the effects of the observer orientation with respect to the recoil velocity vector to take into account the probabilities to measure a given redshifted (or blueshifted) radial velocity of accretion disks with respect to host galaxies. ",
        "introduction": "\\label{sec:Introduction}  Astrophysical black-hole (BH) binaries are characterized by the mass ratio $q=m_1/m_2 \\leq 1$ of the smaller to larger BH, where $m_i$ is the mass of BH $i$, the total mass $M=m_1+m_2$, eccentricity $e$ (assumed to be very small), and spins $\\vec S_i$ (where $S_i/m_i^2 < 1$). In addition, it is often convenient to parameterize the binary with the symmetric mass ratio $\\eta = q/(1+q)^2$, specific spins $\\vec \\alpha_i = \\vec S_i/m_i^2$, total spin $\\vec S = \\vec S_1 + \\vec S_2$, $\\vec \\Delta = M (\\vec S_2 /m_2 - \\vec S_1/m_1)$, and orbital angular momentum $\\vec L$.  The relevance of the spins to the dynamics of black-hole (BH) mergers was recognized soon after the breakthrough in numerical relativity \\cite{Pretorius:2005gq, Campanelli:2005dd, Baker:2005vv} that allowed for the long term, stable numerical evolution of such systems. Notable examples of the early findings are the `hangup' effect \\cite{Campanelli:2006uy}, a repulsive spin-orbit interaction, that delays the merger of black-hole binaries (BHB) when the spins are aligned with the orbital angular momentum, and simultaneously causes the system to radiate excess angular momentum, leading to a remnant BH with sub-maximal spin.  The same mechanism produces an additional attractive effect when the spins are counter-aligned with the orbital angular momentum, leading to a prompt merger. Thus the radiation of angular momentum and energy is asymmetric with respect to the relative orientations of the total spin angular momentum vector and the orbital angular momentum.  When the spins are not exactly aligned or counter aligned, new effects appear (the hangup effect is still present). Precession of the spins is important dynamically because it cause the orbital plane to strongly precess just prior to merger~\\cite{Campanelli:2006fy}.  The final spin of the merged hole can flip with respect to the directions of the individual ones, mainly due to the addition of the orbital angular momentum \\cite{Campanelli:2006fy}.  While the spin-orbit coupling leads to strong precessional effects near merger, the magnitudes of the spins are not affected to the same degree.  In particular spin-orbit interactions are too weak to induce the binary to corotate (or maintain corotation of an initially corotating binary) at the last stages of the merger because the timescale for the radiation driven inspiral is much smaller than the spin-orbit interaction timescale~\\cite{Campanelli:2006fg}.  Numerous other papers have studied different spin effects, such as the large recoil velocities acquired by the remnant of the merger of two spinning black holes~\\cite{Herrmann:2007ac, Koppitz:2007ev, Campanelli:2007ew, Baker:2007gi, Gonzalez:2007hi, Tichy:2007hk} and long term evolutions of generic BHBs (i.e.\\ unequal mass and unequal, randomly-oriented spins)~\\cite{Campanelli:2008nk, Szilagyi:2009qz} to cite a few of the nearly one hundred papers published on the subject since 2006.  The characterization of the remnant black hole (BH) as the by-product of a generic BH binary (BHB) merger  is of great astrophysical interest as it allows one to model the growth of BHs  during the evolution of the universe and their effect on the dynamical evolution of galactic cores and globular clusters, as well as the collisions of galaxies and stellar size binary systems. Thanks to the recent breakthroughs in Numerical Relativity~\\citep{Pretorius:2005gq, Campanelli:2005dd, Baker:2005vv} one can now precisely compute the masses, spins and recoil velocities of these merged BHBs from fully nonlinear numerical simulations.  The modeling of the remnant black hole using fully-numerical techniques was pioneered by the `Lazarus method' \\citep{Baker:2003ds} for spinning black holes followed by the breakthrough `moving puncture' approach. In Refs.~\\citep{Campanelli:2006uy, Campanelli:2006fg, Campanelli:2006fy} the authors studied BHBs characterized by equal-mass, equal-spin individual BHs, with the spins aligned or counter-aligned with the orbital angular momentum, using fully nonlinear numerical calculations and found a simple {\\em ad hoc} expression relating the final mass and spin of the remnant with the spins of the individual BHs. This scenario was later revisited in \\citep{Rezzolla:2007xa, Rezzolla:2007rd} and the formula for the remnant spin was generalized (by assuming that the angular momentum is only radiated along the orbital axis, and neglecting the energy loss) in \\citep{Rezzolla:2007rz} for arbitrary BH configurations (although in the latest paper of this sequel this condition was removed \\cite{Barausse:2009uz}.) In \\citep{Tichy:2008du} a more general {\\em ad hoc} fitting function was proposed. A more comprehensive approach was proposed in \\citep{Boyle:2007ru}; where a generic Taylor expansion, reduced by the physical symmetries of the problem, was used to fit the existing full numerical simulations. A different approach was presented in \\citep{Buonanno:2007sv} where the particle limit approximation was extended to the equal-mass case and the effects of post-ISCO (Innermost Stable Circular Orbit) gravitational radiation were neglected. This approach was further improved in~\\citep{Kesden:2008ga} by taking binding energies into account.  All of these approaches show a certain degree of agreement with the remnant masses and spins obtained in the few dozen fully nonlinear numerical simulations available, but there remains significant uncertainties concerning their accuracy outside this range of parameters.  In this paper we propose a set of formulae that incorporate the benefits of both approaches in a unified way.  Due to the large astrophysical interest of computing remnant recoil velocities, the modeling of recoil velocities followed an independent path, particularly since the discovery \\citep{Campanelli:2007ew, Campanelli:2007cga} that the spins of the black holes play a crucial role in producing recoils of up to $4000\\ \\KMS$. The importance of modeling the recoil velocities as a function of the astrophysical parameters of the progenitor binary was quickly realized \\citep{Campanelli:2007ew, Baker:2007gi, Campanelli:2007cga}.  The news that the merger of binary black holes can produce recoil velocities up to $4000\\ \\KMS$, and hence allow the remnant to escape from major galaxies, led to numerous theoretical and observational efforts to find traces of this phenomenon. Several studies made predictions of specific observational features of recoiling supermassive black holes in the cores of galaxies in the electromagnetic spectrum \\citep{Haiman:2008zy, Shields:2008va, Lippai:2008fx, Shields:2007ca, Komossa:2008ye, Bonning:2007vt, Loeb:2007wz} from infrared \\citep{Schnittman:2008ez} to X-rays \\citep{Devecchi:2008qy, Fujita:2008ka, Fujita:2008yi} and morphological aspects of the galaxy cores \\citep{Komossa:2008as, Merritt:2008kg, Volonteri:2008gj}.  Notably, there began to appear observations indicating the possibility of detection of such effects \\citep{Komossa:2008qd, Strateva:2008wt,Shields:2009jf}, and although alternative explanations are possible \\citep{Heckman:2008en, Shields:2008kn, Bogdanovic:2008uz,Dotti:2008yb},  there is still the exciting possibility that these observations can lead to the first confirmation of a prediction of General Relativity in the highly-dynamical, strong-field regime.  In our approach to the recoil problem \\citep{Campanelli:2007ew, Campanelli:2007cga} we chose to use post-Newtonian theory as a guide to model the recoil dependence on the physical parameters of the progenitor BHB (See Eqs. (3.31) in \\citep{Kidder:1995zr}), while arguing that only full numerical simulations can produce the correct amplitude of the effect. Bearing this in mind, we proposed an empirical formula for the total recoil velocities (see Eq.\\ (\\ref{eq:Pempirical}) below.) Our heuristic formula describing the recoil velocity of a black-hole binary remnant as a function of the parameters of the individual holes has been theoretically verified in several ways. In~\\cite{Campanelli:2007cga} the $\\cos{\\Theta}$ dependence was established and was confirmed in~\\cite{Brugmann:2007zj} for binaries with different initial separations. In \\cite{Herrmann:2007ex} the decomposition into spin components perpendicular and parallel to the orbital plane was verified, and in~\\cite{Pollney:2007ss} it was found that the quadratic-in-spin corrections to the in-plane recoil velocity are less than $20\\ \\KMS$. Recently in \\citep{Lousto:2008dn} we confirmed the leading $\\eta^2$ (where $\\eta = \\frac{m_1 m_2}{(m_1+m_2)^2}$ is the symmetric mass ratio) dependence of the large recoils out of the orbital plane.  Since the magnitude (and direction) of the recoil velocity of the remnant black holes depend so sensitively on the spin orientation just around the time of the formation of a common event horizon, it is important to establish that random oriented spins of individual black holes at large separations (as a plausible initial astrophysical scenario) lead to randomly oriented black holes near merger, or if there is some bias in their orientations by the time they get very close together (i.e.\\ at typical numerical simulations separations of a few $M$, where $M$ is the total mass). It has been argued recently \\cite{Bogdanovic:2007hp,Perego:2009cw,Dotti:2009vz} that the presence of gas and accretion of the individual black holes during the inspiral phase for long time scales can lead to a preferential alignment of spins with the orbital angular momentum, and hence to a configuration that leads to modest recoil velocities (a few hundred $\\KMS$). The latest results by Dotti et al \\cite{Dotti:2009vz} indicate that spin alignment occur in the scale of a few million years within 10 degrees of the orbital angular momentum for cold disks and 30 degrees for warmer disks. The proportion of wet to dry mergers in the universe still needs to be established. Current rough estimates give comparable percentages for both kinds of mergers. With 3 to 5 mergers per galaxy during their lifetime for current spiral and elliptic galaxies respectively \\footnote{M.Volonteri, private communication.}. While the verification of these claims for `wet' mergers is underway, in this paper we would like to explore the possibility that such alignment (or counter-alignment) mechanism exists for purely gravitational interactions (`dry mergers'). In general we will seek to find if there is any bias in the individual spin distributions of black holes at close separations ($5-8M$ for starting typical full numerical evolutions) when starting evolutions with post-Newtonian methods at large  radii with  random spin orientations. These distributions, in turn, will then help in choosing configuration for full numerical simulations of close binaries.   The paper is organized as follows. In Section \\ref{Sec:Inspiral} we describe the post-Newtonian formalism to analyze the inspiral stage of the binary evolutions. We use the Hamiltonian formulation (up to 3.5PN order) to derive the equations of motion in the ADM-TT gauge. Conservative and radiative effects of the spins are included up to the next leading PN order. We also include a purely analytic analysis of the projection of the quantity $\\vec\\Delta$ along the orbital angular momentum $\\vec L$, which has a strong effect on the  recoil velocity, to qualitatively predict a slight bias towards counter-alignment of these two vectors. The results of the statistics of numerical integration of the post-Newtonian equations of motion (EOM) follows. We performed integrations from initial separations of $r=50M$ with $16^4$ spin orientation chosen at random and magnitudes fixed at large astrophysical values, i.e. $S_i/m_i^2=0.97$ for different mass ratios in the range $1/16 \\leq q=m_1/m_2\\leq 1$. The results quantitatively confirm the bias towards counter-alignment of $\\vec\\Delta$ and total spin $\\vec S$ with respect to the orbital angular momentum $\\vec L$. Section \\ref{Sec:Merger} deals with the merger phase, when the black holes are much closer to each other and in a few orbits will merge into a single larger one. This is the typical scenario that full numerical simulations assume. The bulk properties of the remnant black hole can be summarized in terms of empirical remnant formulae that describe its total mass, spin and recoil velocity. We proposed formulae for these quantities based on post-Newtonian scaling with amplitudes fixed by the full numerical simulations. With these formulae at hand, we perform statistical studies by evaluation of these expression for random distributions of mass ratios and individual spins (magnitudes and spins). Those evaluations lead to a large recoil velocity tail in the distribution with non negligible probabilities for $v>1000\\ \\kms$, and highly peaked about the direction of orbital angular momentum at merger. Likewise, evaluations of the final spin formulae lead to a wide distribution peaked at magnitudes of $S_f/M_f^2\\approx0.73$ and orientations peaked at an angle $\\sim25^o$ with respect to the orbital angular momentum. We complete the paper with a discussion of the astrophysical consequences of these results and include an appendix with the computation of the innermost stable circular orbit radius, energy and angular momentum around Kerr black holes (needed for the remnant formulae), with an analytic solution for the equatorial and polar orbits. ",
        "conclusions": "\\label{sec:Discussion}  In this paper we studied the `dry' (i.e. gravitational radiation driven) inspiral of black-hole binaries. We performed the evolutions from separations of the order of $50M$, using up to 3.5PN accurate expressions that represents an excellent approximation for this regime (similar simulations with full numerical relativity could prove practically impossible with current technologies). The statistical results show a small bias towards counter-alignment of the vectors $\\vec\\Delta$ and $\\vec{S}$ with respect to the orbital angular momentum $\\vec{L}$ just prior to merger. This effect essentially takes place at close separations and can be studied analytically at low post-Newtonian orders. The antialignment effect is associated with the late-time precession of the orbital plane due to radiation reaction. This effect for `dry' mergers seems to oppose the alignment mechanism observed in `wet' mergers \\cite{Bogdanovic:2007hp,Perego:2009cw}.  After the initial inspiral regime, we studied the merger of black-hole binaries using full numerical simulations. Here we provided a framework to describe the bulk properties of the remnant of a BHB merger based on PN scaling with free parameters fixed by fitting the results  of full numerical simulations. We have shown how to determine the mass loss in each BHB encounter (\\ref{Eempirical}) and the spin of the remnant BH  (\\ref{Jempirical}) with fitting constants set to match currently available runs. Using the same methods, one can improve the fitting parameters in the above formulae as the results of new runs are made available, thus providing a standard method to incorporate all full numerical results into astrophysical modelings. Precessing, highly spinning binaries, including those with small mass ratios, simulations can provide information to also fix the nonleading terms in the fitting formulae. However, we expect that these extra terms are relatively small and will not have a significant impact on the results presented here.  The new formulae are physically motivated, as they are derived using the post-Newtonian behavior, and naturally incorporate the correct mass ratio $q$ dependence and physical symmetries, as well as allow for the  radiation of angular momentum in the orbital plane. These formulae  model the final plunge of comparable masses black holes in an impulsive approximation are supplemented by the slow inspiral losses that precede this regime by adding the ISCO energy and angular momentum in the particle limit, generalized to symmetric dependence on the mass ratio and spins (Eqs.\\ (\\ref{EISCO}) and (\\ref{JISCO})).  We also extended the  successful recoil formula by adding nonleading terms to include all the linear dependence with the spins, as well as  higher mass ratio powers in Eq.~(\\ref{eq:Pempirical}).  Unlike in the formula for the remnant recoil case, the energy and angular momentum lost by the binary during the inspiral phase is a non-trivial fraction of the total radiated energy and angular momentum (and, in fact, is the dominant contribution in the small mass ratio limit). We thus included both the instantaneous radiative terms for the plunge phase as well as the binding energy and angular momentum at the ISCO into our empirical formulae (\\ref{Eempirical}) and (\\ref{Jempirical}).  Using the fitted coefficients in the above formulae, we find that for equal-mass, non-spinning binaries, the net energy radiated is $5\\%$ of the total mass and the final spin is $\\alpha\\approx0.69$, both in good agreement with the most accurate full numerical runs \\citep{Scheel:2008rj}. For maximally spinning BHBs with spin aligned and counter-aligned  we estimate that quadratic corrections lead to radiated energies between $10\\%$ and $3\\%$ respectively. As for the magnitude of the remnant spin, the linear estimates are between $0.97$ and $0.41$ respectively, with quadratic corrections slightly reducing those values. These results show that the cosmic censorship hypothesis is obeyed (i.e.\\ no naked singularities are formed) and are in good agreement with earlier estimates \\citep{Campanelli:2006fy}.   The set of formulae  (\\ref{Eempirical}) and  (\\ref{Jempirical}) with the fitting constants determined as in the Sec.~\\ref{Sec:Merger} can be used to describe the final stage of binary black holes mergers in theoretical, N-body, statistical studies in astrophysics and cosmology \\citep{O'Leary:2008mq, Blecha:2008mg, Miller:2008yw, Kornreich:2008ca, Volonteri:2007dx, Gualandris:2007nm, HolleyBockelmann:2007eh, Guedes:2008he, Volonteri:2007ax, Schnittman:2007nb, Bogdanovic:2007hp, Schnittman:2007sn, Berti:2008af} by choosing a distribution of the initial intrinsic physical parameters of the binaries $(q, \\vec{S}_1, \\vec{S}_2)$ and mapping them to the final distribution of recoil velocities, spins and masses after the mergers. Here we performed initial studies and have found that: i) The merged black holes have a considerable probability ($23\\%$) to reach recoil velocities above $1000\\ \\KMS$ (See Fig.~\\ref{fig:vfit} and Table~\\ref{table:vdist}) and the distribution is highly peaked along the orbital angular momentum (See Fig.~\\ref{fig:vtheta}). ii) The direction of the spin of the final merged black holes is strongly peaked at an angle of $\\approx25^\\circ$ with respect to the orbital angular momentum pre-merger (see Fig.~\\ref{fig:stheta}), and the spin magnitude is strongly peaked at $S_f/M_f^2\\approx0.73$ (see Fig.~\\ref{fig:sdist}). Higher spins are likely if we include the effects of accretion. This information can be useful in modeling the observational effects of supermassive black holes kicked out of their host galaxies~\\cite{Bogdanovic:2009xd}. For example, as a first approximation, one may assume that the inner accretion disk is associated with the orbital plane of the merging binary, while the direction of the final spin is associated  with the current direction of the radio-jet, and finally that the preferred direction of the kick is along the orbital angular momentum. We can then to reconstruct 3D recoil velocities out of the observer (redshift velocities) information. Also, when modeling the effects of kicks on accretion disks surroundings the merged binary, one should take into account that the most likely recoil velocity depends on the angle with respect to the binary's orbital plane (See Fig. \\ref{fig:v_d_theta}).  In order to take into account that one of the main methods to search for recoiling black holes is to look for large differential redshift, typically between narrow band emission lines coming from the host galaxy and broad emission lines from the portion of the accretion disk that the recoiling black hole carries with it \\cite{Komossa:2008qd, Strateva:2008wt,Shields:2009jf}, we have computed the effect of projecting the computed recoil velocity from the merger of two black holes along the line of sight of an observer on earth. The results are plotted in Fig. \\ref{fig:obs_v_dist} and in Table \\ref{table:vdist}.  Finally, we note that the recoil distributions are sensitive to the assumed distribution of mass ratios. Ideally one should use a distribution consistent with the true distribution of mass ratios of merging galaxies. In~\\cite{Volonteri:2008gj} this distribution is derived analytically from merger scenarios in a cosmological model. They find that the distribution is nearly flat in $\\log_{10} q$ from $q=1/100$ to $q=1$ (see Ref.~\\cite{Volonteri:2008gj}, Fig 1.). When we use a distribution uniform in $\\log_{10} q$ from $-2\\leq \\log_{10}q\\leq0$ here, as suggested in~\\cite{Volonteri:2008gj}, the expected recoil velocity distribution is skewed towards much lower velocities. In Fig.~\\ref{fig:logqdistkick} we plot the recoil magnitude distribution for this distribution, while in Table~\\ref{table:logqvdist} we give the probabilities for obtaining large recoil velocities with this mass ratio distribution. Note that this distribution is, in fact, strongly skewed towards lower mass ratios when compared to the distribution uniform in $q$. Consequently we see much lower probabilities for large recoils. \\begin{figure} \\caption{The distribution of recoil velocity magnitudes and directions (with respect to the orbital plane) assuming a distribution in mass ratios uniform in $\\log_{10} q$ in the interval $-2\\leq\\log_{10}q\\leq0$.} \\includegraphics[width=3.5in]{logqvmag.pdf} \\includegraphics[width=3.5in]{logqvang.pdf} \\label{fig:logqdistkick} \\end{figure} \\begin{table} \\caption{The probability to obtain large recoil velocities, and large recoil velocities along the line of sight if the distribution in mass ratios is uniform in $\\log_{10} q$ for $-2\\leq\\log_{10}q\\leq 0$.} \\label{table:logqvdist} \\begin{ruledtabular} \\begin{tabular}{lllll} $v[\\KMS] \\geq$ & 500 & 1000 & 2000 & 2500\\\\ \\hline Recoil & $20.95\\%$ & $7.99\\%$ & $0.60\\%$ & $0.06\\%$\\\\ Observer & $8.62\\%$ & $2.05\\%$ & $0.06\\%$ & $0.003\\%$ \\end{tabular} \\end{ruledtabular} \\end{table}"
    },
    "0910/0910.2587_arXiv.txt": {
        "abstract": "{ \\vspace{3mm} Invisible $\\psi$ and $\\Upsilon$ decays into light neutralinos, within the MSSM or N(n)MSSM, are smaller than for $\\nu\\bar\\nu$ production, even if light spin-0 particles are coupled to quarks and neutralinos. In a more general way, light dark matter particles are normally forbidden, unless they can annihilate sufficiently through a new interaction stronger than weak interactions (at lower energies), as induced by a light spin-1 $\\,U$ boson, or heavy-fermion exchanges in the case of scalar dark matter. We discuss the possible contributions of $\\,U$-boson, heavy-fermion, or spin-0 exchanges to invisible $\\psi$ and $\\Upsilon$ decays. \\vspace{.5mm} \\\\ \\hspace*{2mm} $U$-exchanges could lead, but not necessarily, to significant branching fractions for invisible decays into light dark matter.\\,\\,We show how one can get the correct relic density together with sufficiently small invisible branching fractions, and the resulting constraints on the $\\,U$ couplings to ordinary particles and dark matter, in particular $\\,|\\,c_\\chi\\,f_{bV}\\,|<\\,5\\ 10^{-3}\\,$ from $\\Upsilon$ decays, for $2\\,m_\\chi$ smaller than a few GeV. We also explain why there is no model-independent way to predict $\\psi$ and $\\Upsilon$ branching fractions into light dark matter, from dark matter annihilation cross sections at freeze-out time. \\vspace{3.5mm}\\\\ PACS numbers: \\,12.60.-i \\,12.60.Cn \\,13.20.Gd \\,14.70.Pw \\,14.80.-j \\,95.35.+d \\hfill \\hbox{LPTENS/09-33} \\vspace{0mm}\\\\ } ",
        "introduction": "\\boldmath New particles in invisible $\\,\\psi\\,$ and $\\,\\Upsilon\\,$ decays.}  \\label{sec:intro}  \\vspace{-1mm}  The nature of dark matter is one of the most challenging issues facing physics. Observation of standard model (SM) particles coupling to invisible final states, as searched for recently in $\\Upsilon$ decays \\cite{babarinv}, might provide information on new neutral particles such as photinos or neutralinos and very light gravitinos, and candidate dark matter constituents \\cite{fayetpsi,fayetkaplan,prd06}. In the standard model, invisible decays of the $\\Upsilon(1S)$ proceed by $b\\,\\bar b$ annihilation into a $\\,\\nu\\bar \\nu$ pair, with a branching fraction~\\cite{fayetkaplan} \\be \\label{upsnu0} B\\,(\\Upsilon(1S)\\,\\to\\,\\nu\\,\\bar\\nu)\\ \\simeq \\ 10^{-5}\\ \\ , \\ee well below the current experimental sensitivity \\cite{babarinv}. However, low-mass dark matter candidates could couple through {\\it stronger-than-weak\\,} interactions to SM particles, and possibly enhance the invisible branching fraction  of the $\\Upsilon(1S)$ to the level of $\\,\\approx 10^{-5}$ to $\\,10^{-2}$ \\cite{fayetkaplan,prd06}, in contrast with weakly-interacting particles, as indicated  by (\\ref{upsnu0}). A new light boson $U$ associated with the gauging of an extra-$U(1)$ symmetry, as considered long ago in \\cite{U}, may play a crucial role as a mediator of such a new interaction.  \\vspace{2mm}  Upper limits on the invisible $\\Upsilon$ branching fraction have been obtained long ago by CLEO  \\cite{upsinv} and ARGUS \\cite{argus}, already having in mind the search for new weakly-interacting particles such as photinos and very light gravitinos, as in invisible $\\psi$ decays \\cite{fayetpsi}. These limits ($7\\ 10^{-3}$ for $\\,\\psi$, $\\,5\\ 10^{-2}$ then $\\,2.3\\ 10^{-2}$ for $\\,\\Upsilon$) are obtained by looking for \\be \\left\\{\\ \\ba{ccc} \\psi\\,(2S) &\\rightarrow& \\ \\pi^+\\,\\pi^-\\ \\psi\\,(1S)_{\\ \\hookrightarrow\\  \\hbox{\\footnotesize invisible}}\\ \\ , \\vspace{2mm}\\\\ \\Upsilon(3S),\\ \\hbox{or}\\ \\Upsilon(2S) &\\rightarrow& \\ \\pi^+\\,\\pi^-\\ \\Upsilon(1S)_{\\ \\hookrightarrow\\  \\hbox{\\footnotesize invisible}}\\ \\ , \\ea \\right. \\ee which provide signatures for the production and invisible decays of $\\,\\psi$ and $\\,\\Upsilon$.    \\vspace{1.5mm}    The $\\Upsilon$ bounds have been improved by Belle and CLEO \\cite{upsinv2,cleo2}, and recently B{\\small A}B{\\small A}R \\cite{babarinv}, down to \\be \\label{limups} B\\,(\\,\\Upsilon(1S) \\ \\to\\ \\hbox{invisible}\\,)\\ < \\ 3\\ 10^{-4}\\ \\ , \\ee at the 90\\,\\% c.l.\\,. We also have, from BES II \\cite{bes}, \\be \\label{limpsi} B\\,(\\,\\psi(1S) \\ \\to\\ \\hbox{invisible}\\,)\\ < \\ 7.2\\ 10^{-4}\\ \\ . \\ee Although the present paper is in general formulated with the $\\Upsilon$, the analysis applies to invisible $\\psi\\,$ decays as well.  \\vspace{1.5mm}   What can we learn about the light neutral particles that could be produced\\,? We shall discuss possible invisible decays of the $\\Upsilon$, \\,and at first  into {\\it light neutralinos} within the MSSM or N(n)MSSM (cf.~Sec.\\,\\ref{sec:ssm}). They are significantly smaller than for $\\Upsilon(1S)$ $\\to\\,\\nu\\bar\\nu$, \\,even in the presence of light spin-0 particles coupling directly quarks to neutralinos.      \\vspace{1.5mm} We shall also discuss, in a more general way, invisible decays of $\\,\\psi\\,$ and $\\Upsilon\\,$ into {\\it \\,light dark matter particles} (of mass $<m_\\psi/2$ \\,or $\\, m_\\Upsilon/2$) \\cite{fayetkaplan,prd06}, which could be much {\\it more strongly coupled to ordinary particles than through ordinary weak interactions}, possibly leading to significant invisible $\\,\\psi$ and $\\,\\Upsilon\\,$ branching fractions.   \\vspace{2mm} Indeed, light dark matter (LDM) particles \\cite{bf,fermion} are normally required to annihilate through a new interaction stronger than weak interactions (at least at lower energies), otherwise their relic density would be too large. Is this compatible with the new experimental bounds  \\cite{babarinv,bes} on  invisible $\\psi$ and $\\Upsilon$ decays\\,? This is part of the more general question \\cite{prd06,fermion,prd07}: how can we have a new interaction stronger than weak interactions, responsible for sufficient annihilations of LDM particles in the early Universe, and at the same time how could it remain unnoticed if it is stronger than weak interactions\\,? Indeed this stronger-than-weak feature cannot persist up to high energies, especially with (production, annihilation or interaction) cross sections increasing like $s$, without getting in conflict with experimental results.  \\vspace{1.75mm}  The apparent contradiction is solved for an interaction mediated by {\\it a new light neutral boson} with small couplings to quarks and leptons, such as the light \\hbox{spin-1} $U$ boson \\cite{U} introduced and discussed long ago, associated with the gauging of an extra-$U(1)$ symmetry. Another possibility, for spin-0 dark matter particles, is obtained for interactions mediated by {\\it heavy-fermion exchanges} \\cite{bf}, which may also allow for sufficient annihilations of light dark matter in the early Universe. We shall also discuss (in Sec.\\,\\ref{sec:ssm}) light spin-0 exchanges, which do not contribute to invisible $\\,\\psi$ and $\\Upsilon$ decays, both in the N(n)MSSM and in a more general way.  \\vspace{1.75mm} The choice is thus, for the production of light dark matter particles in invisible $\\,\\psi$ and $\\Upsilon$ decays,  between \\be \\hbox{\\it a new neutral current} \\ \\ \\  \\hbox{and\\ \\ \\it \\,new heavy fermions}\\,, \\ee or both (reminding us of the early days of gauge theories, before the discovery of the weak neutral current coupled to the $Z$). We shall discuss in Sec.~\\ref{sec:scal} the possible production of scalar dark matter through heavy-fermion exchanges, according to  \\vspace{-7mm} \\be \\Upsilon\\ \\ \\stackrel{\\ba{c}\\hbox{\\small $b_M$}\\vspace{0mm}\\\\ \\ea}{\\longrightarrow}\\ \\ \\varphi\\,\\bar\\varphi\\ \\ . \\ee We shall then concentrate, in Secs.~\\ref{sec:U} and \\ref{sec:comp}, on $U$-in\\-duced reactions, discussing the implications of the experimental limit (\\ref{limups}) on the couplings of the light $\\,U$ boson that may be responsible for $\\,\\Upsilon\\,$ annihilations into light spin-$\\frac{1}{2}$ \\,($\\chi$)\\, or spin-0 \\,($\\varphi$)\\, dark matter particles: \\be \\Upsilon\\ \\ \\stackrel{\\hbox{\\footnotesize $U$}}{\\rightarrow}\\ \\ \\chi\\,\\chi\\ \\ \\ \\hbox{ (or $\\,\\varphi\\,\\bar\\varphi\\,$ or $\\,\\chi\\,\\bar\\chi$)}\\ \\ , \\ee at $\\,\\sqrt s = m_\\Upsilon$ \\cite{prd06}. This $\\,U$ allows for the correct relic density of light dark matter particles \\cite{bf,fermion}, by inducing sufficient annihilations in the early Universe, \\be \\chi\\,\\chi\\ \\ \\ \\hbox{(or $\\,\\varphi\\,\\bar\\varphi\\,$ or $\\,\\chi\\,\\bar\\chi$)} \\ \\ \\stackrel{\\hbox{\\footnotesize $U$}}{\\rightarrow}\\ \\ f\\,\\bar f\\ \\ , \\ee at lower values of the energy, $\\,\\sqrt s\\,\\simeq\\,2\\,m_\\chi\\,$ (or $2\\,m_\\varphi$).  \\vspace{1.75mm}  Finally we shall discuss in Sec.\\,\\ref{sec:rel} whether it makes sense to attempt predicting, in a model-independent way, $\\,\\psi\\,$ and $\\,\\Upsilon\\,$ invisible branching fractions into dark matter, from dark-matter annihilation cross sections at freeze-out time; and consider briefly, in Sec.\\,\\ref{sec:npert}, possible non-perturbative effects associated with light $U$ exchanges.      \\vspace{-.5mm} ",
        "conclusions": "\\vspace{-1mm}   Searching for invisible $\\psi$ and $\\Upsilon\\,$ decays has long been identified and used as a way to search for new light particles, especially if they are more strongly coupled to standard model ones than through weak interactions. Invisible branching fractions into neutrinos in the SM and neutralinos in the (N/n)MSSM, however, are well below present experimental limits, even if light spin-0 particles are directly coupled to both quarks and neutralinos.  \\vspace{1mm} Exchanges of a new heavy quark $\\,b_M$ or $c_M$ also lead to small invisible branching fractions into scalar dark matter, unless the $\\,b_M$ or $c_M$ quarks are only moderately heavy, and the corresponding Yukawa couplings to scalar dark matter rather large.  \\vspace{1mm}  Exchanges of a light spin-1 $U$ boson associated with a new gauge interaction, with a vector coupling to quarks, may induce invisible decays of $\\,\\psi$ and $\\,\\Upsilon$ into light dark matter particles at a significant rate, or, conversely, at a very small rate, with $\\,U$ exchanges responsible for sufficient dark matter annihilations in the early Universe.   \\vspace{.5mm} The new limit on invisible $\\Upsilon\\,$ decays constrain the $U$ couplings to dark matter and $b$ quarks to satisfy $\\,\\,|\\,c_\\chi\\,f_{bV}\\,|\\,<\\,5\\ 10^{-3}$, \\,for $\\,2\\,m_\\chi\\,$  smaller than a few GeV's. This may be compared with $\\,\\,|\\,c_\\chi\\,f_{cV}\\,|\\,<\\,9.5\\ 10^{-3}$ from BES II (expected to get significantly improved at BES III). Small values of $m_\\chi$ of a few MeV's, and  $\\,m_U$ of less than a few hundred MeV's, in fact, may well be preferred.   \\vspace{.5mm}  The radiative decays $\\,\\psi\\,$ or $\\Upsilon\\to\\gamma\\,\\chi\\chi\\,$, or $\\,\\gamma\\,\\varphi\\,\\bar\\varphi\\,$, may be induced by the axial couplings of the $U$ to the $c$ and $b$ quarks, with amplitudes proportional to $\\,c_\\chi f_{cA}\\,e$ or $\\,c_\\chi f_{bA}\\,e$. If they can be sufficiently well constrained although the emitted photon is not monochromatic, they may lead to interesting limits on $\\,c_\\chi\\,f_{cA}$ or $\\,c_\\chi\\,f_{bA}$, \\,especially for larger $\\,c_\\chi$, as $f_{cA}$ and $f_{bA}$ are already constrained from $\\,\\psi\\,$ or $\\Upsilon\\to\\gamma\\,U$. These decays may also be sensitive to heavy-quark exchanges, for scalar dark matter.  \\vspace{.5mm}   Even if standard decay modes into neutrinos are not reachable yet, searches for invisible meson decays give useful informations on the $U$ boson and its couplings, contributing to shed light on the nature of dark matter and on the dark force through which it interacts with ordinary particles."
    },
    "0910/0910.0850_arXiv.txt": {
        "abstract": "In this paper we consider the effects of opacity regimes on the stability of self-gravitating protoplanetary discs to fragmentation into bound objects. Using a self-consistent 1-D viscous disc model, we show that the ratio of local cooling to dynamical timescales $\\Omega \\tc$ has a strong dependence on the local temperature.  We investigate the effects of temperature-dependent cooling functions on the disc gravitational stability through controlled numerical experiments using an SPH code.  We find that such cooling functions raise the susceptibility of discs to fragmentation through the influence of temperature perturbations -- the average value of $\\Omega \\tc$ has to increase to prevent local variability leading to collapse. We find the effects of temperature-dependence to be most significant in the `opacity gap' associated with dust sublimation, where the average value of $\\Omega \\tc$ at fragmentation is increased by over an order of magnitude.  We then use this result to predict where protoplanetary discs will fragment into bound objects, in terms of radius and accretion rate.  We find that without temperature dependence, for radii $\\lesssim 10 AU$ a very large accretion rate $\\sim 10^{-3} \\msolaryr$ is required for fragmentation, but that this is reduced to $10^{-4} \\msolaryr$ with temperature-dependent cooling.  We also find that the stability of discs with accretion rates $\\lesssim 10^{-7} \\msolaryr$  at radii $\\gtrsim 50 AU$ is enhanced by a lower background temperature if the disc becomes optically thin. ",
        "introduction": "\\label{intro}  The formation of planets within protoplanetary discs is a subject that attracts considerable interest, with two main competing schools of thought. The core accretion-gas capture model \\citep{Lissauer93, LissauerS07, Klahr08} posits  hierarchical growth, with the collisional coagulation of dust grains initially leading to centimetre-sized particles, and thence on to planetesimals and rocky planets.  Once a critical mass is reached, it is then possible to accrete a gaseous envelope and hence form giant Jupiter-like planets.  Various observations have successfully confirmed this mode of planet formation, for example \\citet{Marcyetal05}, \\citet{Dodson-RobinsonB09}.  However, this model cannot explain all the available observations. \\citet{KennedyK08} show that beyond approximately 20 AU, the timescales for giant planet formation via core accretion exceed the expected disc lifetime of approximately 10 Myr, implying that no planets should be detected in this region.  However, recent observations of HR8799 with the Keck and Gemini telescopes have produced direct images of giant planets (5 - 13 $M_{\\mathrm{J}}$) orbiting at radii of up to $\\sim 70$ AU \\citep{Maroisetal08}. Similar observations of other systems (e.g. $\\beta$ Pic b \\citep{Lagrangeetal09} and Formalhuat \\citep{Kalasetal08}) and theoretical work on the formation of 2MASS1207b \\citep{LodatoDC05} have suggested that there is another mechanism for planet formation at work, and this is thought to be the effect of gravitational instabilities within the protoplanetary discs themselves.  In protoplanetary discs where the self-gravity of the gas is dynamically important, direct gravitational collapse of locally Jeans-unstable over-densities within the disc \\citep{Boss97, Boss98, Durisenprpl} would also produce giant planets very rapidly, on the local dynamical timescale.  A similar process of gravitational instability leading to local collapse is a strong candidate for the formation of stellar discs around Active Galactic Nuclei (AGN) \\citep{NayakshinCS07} and those observed in our own Galactic Centre \\citep{LevinB03, NayakshinC05}, and in the context of protostellar discs may furthermore be responsible for the formation of brown dwarves and other low mass stellar companions \\citep{StamatellosHW07}.  The emergence of the gravitational instability within a disc is governed by the parameter $Q$ \\citep{Toomre64}, which for a gaseous Keplerian disc is given by \\begin{equation} Q = \\frac{\\cs \\Omega}{\\pi G \\Sigma}. \\label{Q} \\end{equation} This encapsulates the balance between the stabilising effects of rotation ($\\Omega(R)$ is the angular frequency at radius $R$) and thermal pressure ($\\cs (R)$ is the sound speed) and the destabilising effect of the disc self-gravity via the surface density $\\Sigma(R)$.  When $Q \\lesssim 1$ the instability is initiated, leading to the presence of spiral density waves within the disc which, depending on the local cooling rate, may persist in a self-regulated quasi-stable state \\citep{Gammie01, LodatoR04} or may fragment into bound clumps \\citep{JandG03,   RiceLA05}.  Discs that are sufficiently cool are therefore expected to be susceptible to the gravitational instability, and it is thought that at least in the early stages of stellar evolution, many discs enter the self-gravitating phase \\citep{Hartmannbook}.  Once the gravitational instability is initiated, heat is input to the disc on the dynamical timescale through the passage of spiral compression/shock waves \\citep{Cossinsetal09}.  Various numerical studies using both 2D and 3D models of self-gravitating discs have produced the result that, in order to induce fragmentation, the disc must be able to cool on a timescale faster than a few times the local dynamical time, $t_{\\mathrm{dyn}} = \\Omega^{-1}$ \\citep{Gammie01, RiceLA05}.  This condition is likely to occur only at relatively large radii ($\\sim 100$ AU), on the assumption that stellar or external irradiation of the disc is negligible \\citep{Rafikov09, StamatellosW09}.  These models have generally used a cooling rate prescribed by using a fixed ratio between the local cooling ($\\tc$) and dynamical ($\\Omega^{-1}$) times, such that \\begin{equation} \\Omega \\tc = \\beta \\end{equation} for some constant $\\beta$ throughout the radial extent of the disc.  Various authors (e.g. \\citealt{Gammie01, RiceLA05}) have found that fragmentation occurs whenever $\\Omega \\tc \\approx 3 - 7$.  By using a more realistic cooling framework based on the optical depth, \\citet{JandG03} found that the fragmentation boundary (defined hereafter as the ratio of the cooling to dynamical timescales, $\\Omega \\tc$ at fragmentation) may in fact be over an order of magnitude greater than this, leading to an enhanced tendency towards fragmentation.  This variation in $\\Omega \\tc$ they ascribed to the implicit dependence of the cooling function on the disc opacity, and hence on temperature.  Using the opacity tables of \\citet{BellLin94} it is clear that the opacity is a strong function of temperature in certain regimes, and by modelling protoplanetary discs as optically thick in the Rosseland mean sense, $\\Omega \\tc$ shows power law dependencies on both the local temperature and density. In cases where this dependence is strong, it is therefore possible that small temperature fluctuations may push the local value of $\\Omega \\tc$ below the fragmentation boundary, even when the \\textit{average} value is significantly above it.  In this paper we therefore seek to investigate and clarify the exact relationship between the fragmentation boundary and the temperature dependence of $\\Omega \\tc$, using a Smoothed Particle Hydrodynamics (\\textsc{sph}) code to conduct global, 3D numerical simulations of discs where the cooling time follows a power-law dependence on the local temperature.  In addition, various studies have shown that in a quasi-steady state the gravitational instability may be modelled pseudo-viscously (\\citealt{LodatoR05, Cossinsetal09, Clarke09, Rafikov09}).  We therefore use the $\\alpha$-prescription of \\citet{ShakSun73} and the assumption of local thermal equilibrium, where \\begin{equation} \\Omega \\tc = \\frac{4}{9} \\frac{1}{\\gamma (\\gamma - 1) \\alpha} \\label{thermalbalance} \\end{equation} and $\\gamma$ is the ratio of specific heats, to construct an analytical model of the opacity regimes present within a marginally gravitationally stable disc.  From this we can therefore predict analytically if and where such discs would become prone to fragmentation, and also compare these results to those from the more complex simulations where radiative transfer is modelled, such as \\citet{Boley09} and \\citet{StamatellosW09}.  The structure of this paper is therefore as follows.  In Section 2 we discuss some of the theoretical results relevant to protoplanetary discs, and introduce a simplified cooling function derived from the various opacity regimes.  We further consider the effects we expect these cooling prescriptions to have on the susceptibility of protoplanetary discs to fragmentation.  In Section 3 we briefly outline the numerical modelling techniques used in our simulations and detail our initial conditions.  In Section 4 we present the results from these simulations, before proceeding to collate these with the analytical predictions in Section 5.  Finally in Section 6 we discuss the ramifications of our work and the conclusions that may be drawn from it. ",
        "conclusions": "\\label{discussion}  In summary, we have found from controlled numerical experiments with an imposed temperature dependent cooling law that the effect of temperature dependence is to increase the value of $\\Omega \\tc$ at which the disc will fragment into bound objects.  Furthermore, this tendency to fragment is greater the more strongly the cooling function depends on the local disc temperature.  In this respect, this confirms the results of \\citet{JandG03}, who likewise noted a markedly increased tendency towards fragmentation in certain opacity regimes.  This result has been attributed to uncertainty in the value of $Q$ in the self-regulated state \\citep{Clarke09}, equivalent to uncertainty in the equilibrium temperature in our models.  However, our results show that this is only one of two mechanisms that affect the fragmentation boundary, and one that we have been able to account for \\textit{a posteriori} by using effective values of $\\beta$ rather than those input to the simulations.  The other effect is due to the strength of the intrinsic temperature perturbations about the mean.  In the case where the cooling law is dependent on temperature, perturbations about the equilibrium temperature will mean that some fraction  of the gas has a \\textit{lower} value of $\\Omega \\tc$ than average.  Once this fraction reaches a critical value, the disc will become unstable to fragmentation.  As the dependence of the cooling on these temperature perturbations increases, at a given average value of $\\Omega \\tc$ the percentage of gas that lies below the critical value also increases, and thus the average must increase to avoid fragmentation.  We therefore find that the effect of allowing the cooling function to depend on the local temperature is to make the disc more unstable to fragmentation, and we have been able to quantify this variation (see \\eref{betan_empirical}). Combining this with predictions of the temperature dependence of protoplanetary discs using opacity-based cooling functions, we find that the fragmentation boundary can be increased by approximately an order of magnitude in terms of $\\Omega \\tc$, in close agreement with \\citet{JandG03}. We have also found that the RMS strength of the temperature perturbations can be correlated to the average cooling strength (see \\eref{kempirical}), in a very similar manner to that found for the surface density fluctuations \\citep{Cossinsetal09}.  Using these predicted values in analytic models of marginally-gravitationally stable $Q = 1$ discs with a representative equation of state, we have found that the susceptibility of such discs to fragmentation into bound objects is also sensitive to the steady state mass accretion rate, as shown in \\fref{Omtc-MdotR}.  Others have noted that in the optically thick limit where the opacity is dominated by ices, $\\Omega \\tc$ is \\textit{independent} of temperature, and thus the cooling rate is determined only by the local density, itself a function of radius \\citep{MatznerL05, Rafikov05, Clarke09}. It has therefore been suggested that once the cooling becomes dominated by ices fragmentation beyond some radius on the order of 100 $AU$ becomes inevitable, and indeed we find that with a background ISM temperature of $10K$, fragmentation occurs at $\\sim 50 AU$ for all accretion rates below $\\sim 10^{-5} \\msolaryr$.  However, if this minimum temperature condition is relaxed, we find that the change in cooling due to entering the optically thin regime has the effect of stabilising the disc out to large radii.  (The fact that allowing it to become cooler actually \\textit{stabilises} the disc is due to the fact that in this regime $\\Omega \\tc$ increases with decreasing temperature, and thus a hot disc has a shorter cooling time than a cold one.)  For Class II / Classic T Tauri objects embedded in a cold medium with accretion rates below a few times $10^{-7} \\msolaryr$, it is therefore possible that extended discs well beyond 100 $AU$ may be stable against fragmentation (they may well be stable against gravitational instabilitites altogether), and indeed discs with radii of at least 200 $AU$ have been observed (see for example \\citealt{Eisneretal08}). Nonetheless, discs with accretion rates at the higher end of the scale ($\\dot{M} \\approx 10^{-6} \\msolaryr$, \\citealt{Hartmannbook}) will still be unstable to fragmentation at radii beyond $\\sim 50 AU$.  It should be borne in mind however that in the outer regions of discs where the surface density is low, non-thermal ionisation (from cosmic rays, X-rays etc) can trigger the MRI, and this may provide an alternative mechanism for preventing fragmentation, as shown in \\citet{Clarke09}.  \\fref{Omtc-MdotR} also shows another important result, that for accretion rates between $10^{-8} - 10^{-2} \\msolaryr$ discs cannot exist in a non-fragmenting purely self-gravitating state at radii $\\lesssim 2 - 5AU$.  In this regime discs are either MRI active ($\\dot{M} \\lesssim 10^{-4} \\msolaryr$) or unstable to fragmentation ($\\dot{M} \\gtrsim 10^{-4} \\msolaryr$).  We also find that in a narrow band of accretion rates $\\sim 10^{-4} \\msolaryr$ it is possible for discs to be both MRI active and unstable to fragmentation, although the exact interaction of these two instabilities is uncertain (see \\citealt{Fromangetal04}).  It is therefore the case that for steady-state protoplanetary discs the gravitational instability cannot drive accretion directly onto the protostar -- either the MRI or the thermal instability must act at low radii, as has been proposed for FU Orionis outbursts (\\citealt{Armitage01, Zhuetal09}).  Finally, our results agree with the generally accepted view that planet formation through gravitationally-induced fragmentation is unlikely to occur at radii less than 50 - 100 $AU$ \\citep{MatznerL05, Rafikov05, WhitworthS06, Clarke09, Rafikov09}, although this critical radius varies with both the mass accretion rate and the background ISM temperature.  Within this radius the core accretion model remains likely to be the dominant mode of planet formation.  Outside this radius however, the fragmentation of spiral arms will produce gaseous planets, a result which matches that of \\citet{Boley09} using a grid-based hydrodynamical model with radiative transfer -- fragmentation was noted at $\\sim 100 AU$ about a $1 M_{\\odot}$ protostar.  This result is further corroborated by \\citet{StamatellosW08} whose radiative transfer SPH code suggested a massive disc about a $0.7 M_{\\odot}$ protostar would rapidly fragment into planetary mass objects or brown dwarf companions beyond approximately $100 AU$. Although the mass accretion rate onto the central object is not stated in either case, we find that these figures are nonetheless in general agreement with our predictions."
    },
    "0910/0910.5191_arXiv.txt": {
        "abstract": "We study the line profiles of a range of transition region (TR) emission lines observed in typical quiet Sun regions. In magnetic network regions, the \\ion{Si}{4}~1402 \\AA{}, \\ion{C}{4}~1548 \\AA{}, \\ion{N}{5}~1238 \\AA{}, \\ion{O}{6}~1031\\AA{}, and \\ion{Ne}{8}~770 \\AA{} spectral lines show significant asymmetry in the blue wing of the emission line profiles. We interpret these high-velocity upflows in the lower and upper TR as the quiet Sun equivalent of the recently discovered upflows in the low corona above plage regions \\citep{Hara2008}. The latter have been shown to be directly associated with high-velocity chromospheric spicules that are (partially) heated to coronal temperatures and play a significant role in supplying the active region corona with hot plasma \\citep{DePontieu2009}. We show that a similar process likely dominates the quiet Sun network. We provide a new interpretation of the observed quiet Sun TR emission in terms of the relentless mass transport between the chromosphere and corona - a mixture of emission from dynamic episodic heating and mass injection into the corona as well as that from the previously filled, slowly cooling, coronal plasma. Analysis of the observed upflow component shows that it carries enough hot plasma to play a significant role in the energy and mass balance of the quiet corona. We determine the temperature dependence of the upflow velocities to constrain the acceleration and heating mechanism that drives these upflows. We also show that the temporal characteristics of these upflows suggest an episodic driver that sometimes leads to quasi-periodic signals. We suggest that at least some of the quasi-periodicities observed with coronal imagers and spectrographs that have previously been interpreted as propagating magnetoacoustic waves, may instead be caused by these upflows. ",
        "introduction": "One of the most overlooked but important problems in solar and stellar physics is the nature of the transport mechanism that is required to supply enough mass from the cool chromosphere to the hot corona to counter the effects of the slow, but steady cooling and draining of the corona on timescales of order half an hour \\citep{Schrijver2000}. Ever since their discovery in the late 1800s when ``vertical flames'' where first seen at the limb during total solar eclipse \\citep[][]{Secchi1877}, spicules or dynamic jet-like extrusions \\citep[][]{Roberts1945,Beckers1968} have received much attention as agents in the transport of energy and mass in the lower solar atmosphere \\citep[e.g.,][]{Thomas1948, Pneuman1978, Athay1982}. This is because the mass they carry upward is estimated to be two orders of magnitude larger than that needed to supply mass to the corona and solar wind \\citep[e.g.,][]{Athay1982}. One shortcoming was that a signature of these ejecta could not be observed at ``typical'' coronal temperatures \\citep[e.g.,][]{Withbroe1983} and so the concept of spicules as a hot mass transport mechanism for the outer solar atmosphere has been largely mothballed.  The Solar Optical Telescope \\citep[SOT;][]{Tsuneta2008} of \\hinode{} \\citep{Hinode} has revolutionized our view of the dynamic chromosphere revealing the existence of at least two types of spicules \\citep[][]{DePontieu2007b}. ``Type-I'' spicules are long-lived (3-5 minutes) and exhibit longitudinal motions of the order of 20km/s that are driven by shocks resulting from the leakage of p-modes and convective motions, as well as magnetic energy release in regions around strong magnetic flux concentrations \\citep[][]{DePontieu2004,Hansteen2006, DePontieu2007a,Martinez-Sykora2009}. Their contribution to TR emission may be limited to a hot sheath with little of the chromospheric plasma heated to TR temperatures \\citep{deWijn2006,Judge2008}. ``Type-II'' spicules, on the other hand, inhabit a truly different dynamic regime than the ``classical'' spicules studied by \\citet[][]{Roberts1945} (or others) and have a direct impact on the TR.  These recently discovered spicules show much shorter lifetimes (50-100s) and higher velocities (50-150~km/s), and are likely caused by reconnection \\citep[][]{DePontieu2007b}.  Type II spicules often exhibit only upward motion, with the feature fading along its entire length in a matter of seconds, which is suggestive of rapid heating to (at least) TR temperatures.  \\citet[][]{DePontieu2009} (hereafter D2009) recently studied active region plage and discovered a direct connection between Type-II spicules and a persistent, faint, but strongly blue-shifted component of emission in hot coronal lines at the magnetic footpoints of loops.  Their work suggests that coronal heating often occurs at chromospheric heights in association with jets that are driven from below and supply the corona with a significant fraction of the required mass flux.  Here we show that a similar process occurs in the quiet Sun for a wide range of temperatures that cover both the lower and upper TR (or low corona). In contrast to D2009, the current paper is focused on the quiet Sun. We study several aspects of this mass-loading process that were not addressed in D2009. In \\S\\ 2 we discuss the quiet Sun observations and present a novel analysis technique for UV emission lines observed by SUMER \\citep[][]{Wilhelm1995} on \\soho{} \\citep[][]{Fleck1995}. We perform a detailed study of the strengths and limitations of the line asymmetry analysis. We investigate the effects of blends on the analysis, and compare data observed in both first and second order to show that our conclusions are not impacted by blends. We describe the temporal behavior of the blueward asymmetries and compare them to those of type~II spicules. Most importantly, we put forth a new interpretation of the nature of the emission from the quiet Sun transition region in terms of the mass cycling associated with the heating and cooling of the plasma associated with the type-II spicules. We discuss this new interpration in \\S\\ 3. ",
        "conclusions": "\\label{discuss} We observe high speed upflows in emission lines that are formed, under equilibrium conditions, in the low, middle, and upper TR (or low corona). The upflows show a similar spatial pattern in all lines, preferentially occur around the network regions, and have velocities of order 40-100 km/s. The emission of the upflowing plasma is faint at $\\sim$ 5\\% of the peak intensity of the line. It is possible that these upflows are visible as propagating disturbances in EUV \\citep[][]{Schrijver1999} and soft X-Ray coronal images \\citep[][]{Sakao2008}: the velocity of the measured weak episodic jets are in the same range as the apparent velocities of the disturbances in imaging data. A connection to the chromospheric activity and/or spectroscopic measurements is studied by \\citep{McIntosh2009b} for active regions.  What causes these upflows? Many authors have discussed blueward excursions of TR lines in the past \\citep[e.g.][]{Dere1989}, with most of those observations being linked to explosive events: violent, often bi-directional jets caused by reconnection \\citep{Chae1998}. What we report on here is not an occasional disturbance, like an explosive event, but a much more ubiquitous and fainter phenomenon: at any single moment, most of the network regions show these high speed upflows. We believe that these upflows are related to the blueward asymmetry and line broadening that has been reported before for TR lines in the quiet Sun network \\citep[e.g.][]{Peter2000}. The latter suggested, without the benefit of the R-B measure, that hot coronal funnels cause this faint component, with the dominant core of the line emitted in small-scale, low-lying coronal loops. Our interpretation is very different. The quiet Sun upflows are very similar in location, velocity (50-120 km/s), intensity (5\\% of line core) and temporal behavior to the faint, coronal upflows observed with Hinode/EIS above active region plage \\citep[][D2009]{Hara2008}. The latter have been directly linked to the heating to coronal temperatures of chromospheric type II spicules, which dominate the upper chromosphere in plage and quiet Sun network and develop velocities of order 50-150 km/s. D2009 estimated that the mass flux injected to the corona by these jets is significant for the mass balance of active region coronal loops. We suggest that the upflows we observe here in \\ion{Si}{4}, \\ion{C}{4}, \\ion{N}{5}, \\ion{O}{6}, and \\ion{Ne}{8} are the quiet Sun equivalent of this process.   In Fig.~\\pref{f15} we provide a cartoon of the TR structure consistent with our observations and the analysis of D2009. What we observe as the transition region is a conveyor belt of sorts, one that includes frequent bursts of rapidly upflowing ionized hot plasma (in the form of thin, dynamic type II spicules), and recycles the material that is slowly cooling out of the corona, losing energy through radiation and producing the well known prevalent TR red-shift, as well as a host of other TR phenomenology discussed above \\citep[e.g.,][]{Gebbie1981, deWijn2006, Judge2008}. This figure is very different from Fig.~8 of \\citet{Peter2000}. We interpret the supergranular magnetic network as a truly dynamic environment of very high speed ionized plasma upflow (in the form of type II spicules) with a mix of instantaneous and gradual mass return from the corona {\\em instead} of a component due to ``hot coronal funnels\" rooted in supergranular boundaries. While low-lying, cooler loops \\citep[e.g.,][]{Dowdy1986} may well play some role\\footnote{\\citet{Judge2008a} recently questioned this view of transition region structure.} in the TR emission \\citep{Peter2000} our analysis has demonstrated that the role of spicules is fundamental in understanding the observed TR signature. It is possible that the small network core blueshifts observed in \\ion{Ne}{8} may be the result of a more gentle evaporative mass loading due to coronal energy deposition \\citep[e.g.,][]{Klimchuk2006,Tomczyk2009} in the upper TR. Regardless of the interpretation, our results suggest that approximating the observed spectral line profiles with (even multiple) Gaussians is profoundly limiting, and that other ways to consistently model the observed asymmetric plasma velocity distributions are essential \\citep[e.g.,][]{Scudder1992a}.  Our conceptual view of the TR is related to the model put forth by \\citet{Pneuman1978} who suggested that a small fraction of spicules are heated to coronal temperatures and provide the corona with hot plasma. To constrain the spicular mass flux into the quiet Sun corona we follow the approach of D2009, appealing to the conservation of mass and observed emission measures in the system and noting that the latter are similar to those observed in active regions - upflow components with $\\sim 5\\%$ of the core intensity. D2009 found that the density of the coronal component of the spicules could be expressed as $\\rho_i \\approx \\rho_c /(3\\epsilon_T^2)$ where $\\rho_c$ the coronal density, and $\\epsilon_T$ is the fraction of total coronal spicular density at a temperature of $T$. If we assume that only a fraction $\\epsilon_T=0.25-0.5$ of the spicular coronal mass flux is observed in \\ion{Ne}{8}, then the only difference with the analysis of D2009 is that the coronal density in quiet Sun is less than that of active regions (say, 3x10$^8$ instead of $10^9$ cm$^{-3}$). As a result the coronal densities $\\rho_c$ of hot spicular matter are of order $3 - 15 \\times 10^8$ cm$^{-3}$. This implies that the heated spicules can fill the quiet Sun corona even if only 0.3-1.5\\% of the chromospheric spicule mass reaches coronal temperatures. D2009 also estimate that the number of spicules per arcsec$^2$ occurring within the network over the lifetime of a spicule is $N_n \\approx 2 \\epsilon_T^2  \\, 725^2/\\delta_i^2$, where $\\delta_i \\approx 250$km the diameter of a typical spicule. Since the assumptions about the coronal scale height, expansion from network to corona, velocity and diameter of spicules are of the same order for active region and quiet Sun, we estimate a density of spicules of order 1-4 per arcsec$^2$ within a network region. Is this number compatible with Hinode/SOT observations of type II spicules at the limb? \\citet{Rouppe2009} report that they see at least between 1 and 3 of these spicules per linear arcsec at the limb (at a height of 5,000 km). \\citet{Rouppe2009} indicate that when observing at 5,000 km height above the limb, line-of-sight superposition implies that one can see features emanating from network regions over a line-of-sight that is of order 200 arcseconds long. Such a line-of-sight in principle could cross network at the edge of of $\\sim 4$ supergranules. However, the network around supergranular cells has a relatively low filling factor of order 0.25, as shown in, e.g., Fig.~12 of \\cite{Rouppe2009}. This means that our estimate of 1-4 spicules per arcsec$^2$ within a network region fits in well with the observed number of spicules (1-3 per linear arcsec) that is observed with Hinode at the limb (since the number of supergranule cells multiplier is cancelled by the low network filling factor). We conclude that the observed number of spicules at the limb is fully compatible with a scenario in which the heated component of type II spicules observed in TR emission can play a significant role in the mass balance of the quiet corona.  \\begin{figure} \\plotone{f15} \\caption{Conceptual view of TR emission in network and inter-network.\\label{f15}} \\end{figure}  The impact of these spicules on the coronal energy balance can be estimated by calculating the flux associated with the internal energy of the mass flux calculated in the above. We find that the energy flux is given by $N_i \\rho_c c_v T v_i$, with $N_i$ the number of spicules that occur at any location over the spicule lifetime, $\\rho_c$ the mass density of the spicule component that is heated to coronal temperatures, $c_v$ the specific heat at constant volume, $v_i$ the upflow velocities in the spicules, and $T$ the temperature the coronal spicular component reaches. D2009 estimated that $N_i \\approx 2 \\epsilon_T^2$. If we assume that all the plasma in these spicules is heated to 1~MK (of order the formation temperature of Ne VIII, this implies $\\epsilon_T \\approx 1$), and use $\\rho_c = 3 \\times 10^8$ cm${-3}$, $v_i \\approx 100 $ km s$^{-1}$, we find an energy flux of order $7 \\times 10^5$ erg cm$^{-2}$ s$^{-1}$. This is of the same order of magnitude as the estimated conductive and radiative losses \\citep[$3 \\times 10^5$ erg cm$^{-2}$ s$^{-1}$ for quiet Sun,][]{Priest1982}.  Several issues are unresolved in the scenario presented here. The first is the role that spicules play in the mass balance of the lower TR. Observations from Hinode/SOT indicate that many are heated above 20,000K \\citep[][]{DePontieu2007b}, but the peak temperature that the bulk of the spicular mass reaches is unclear. \\citet{Pneuman1978} suggest that the fraction of spicular mass reaching a certain TR temperature $T$ is a steeply declining function of $T$, consistent with the rapid decrease of differential emission measure from 10,000 to 100,000K. Simultaneous observations with high spectral, temporal and spatial resolution of the chromosphere and TR at high signal-to-noise ($>$20) are necessary to resolve this issue. With the recent selection of the Interface Region Imaging Spectrometer (IRIS) such data should be forthcoming.  A second issue relates to the formation and heating mechanism of type II spicules themselves. Insight into the magnetic geometry required for the energy release and how the plasma itself is heated and accelerated along the structure of the spicule may help explain why we see a velocity difference in the upward velocity components between the formation conditions of \\ion{Si}{4}, \\ion{C}{4}, \\ion{N}{5}, \\ion{O}{6} and \\ion{Ne}{8} as well as those between the quiet and active atmosphere (with upflows stronger in the latter -- up to 150 km/s). Certainly, the temporal behavior points to a driver that leads to episodic upflows on characteristic timescales of granulation, something also seen by \\citet{Rouppe2009} who report on the discovery of the disk-counterpart of type~II spicules.  Finally, based on the observational evidence and interpretation presented in this paper, efforts to better understand type II spicules may lead to a clearer picture of the impact of non-equilibrium ionization on the upper portion of the solar atmosphere."
    },
    "0910/0910.2778_arXiv.txt": {
        "abstract": "{ In this paper, we describe a very simple method to calculate the positions of the planets in the sky. The technique used enables us to calculate planetary positions to an accuracy of $1^{\\circ}$ for $\\pm 50$ years from the starting epoch. Moreover, this involves very simple calculations and can be done using a calculator. All we need are the initial specifications of planetary orbits for some standard epoch and the time periods of their revolutions. } ",
        "introduction": "The night-sky fascinates people. To be able to locate a planet in the night sky is something that thrills people. Since the planets move with respect to the background stars and continuously change their positions in the sky, locating them in the sky could appear be a non-trivial task. It is a general notion that calculating the planetary positions is a very tedious task, involving a lot of complicated mathematical equations and computer work. However, to be able to locate planets in the sky one does not really need very accurate positions. After all, Kepler's laws, which describe planetary orbits reasonably well, are mathematically simple. Hence, one could use Kepler's law to predict planetary positions in which mutual influence of planets is not considered. Thereby an accuracy of $\\sim 1^{\\circ}$ in planetary positions would be achieved.  In this project, we employ a very simple method to calculate the positions of the planets. The technique we use enables us to calculate planetary positions to an accuracy of $1^{\\circ}$ for $\\pm 50$ years from the starting epoch. Moreover, this involves very simple calculations and can be done using a calculator. All we need are the initial specifications of planetary orbits for some standard epoch and their time periods of revolution. Although accurate planetary positions could be obtained easily from the internet, yet it is very instructive and much more satisfying to be able to calculate these ourselves, starting from basic principles and using a simple procedure.  Our first step would be to calculate the positions of all the planets (including Earth) in their orbits around the Sun. We initially consider the planets to revolve around the Sun in uniform circular motions. Knowing their original positions for the starting epoch, we calculate their approximate positions for the intended epoch. As a consequence of this approximation there will be an error since the actual orbits are elliptical. To get more accurate positions, we require some corrections, which are derived in Appendix A. These corrections account for the elliptical motion.  Knowing the positions of the planets around the Sun, we can then use simple co-ordinate geometry to transform their position with respect to an observer on Earth. Our task becomes simple since the orbits of all planets more or less lie in the same plane, viz. the ecliptic plane.  In this project, we calculate the motion of naked-eye planets only, although the procedure can be applied equally well for the remaining planets also. ",
        "conclusions": "It is a general notion that calculation of position of planets in the night sky is a difficult job, which can be accomplished only by complex scientific computations, using fast computers. The motive of this project has been to bring out the fact that such complex and accurate computations are not always really necessary. One can calculate the position of planets using the method derived here and get the thrill of finding the planet at the predicted position in the night sky.  We have been able to obtain the position of planets within an accuracy of $1^{\\circ}$, using a calculator. This method can be used to reckon planetary positions up to $\\pm 50$ years of the starting epoch."
    },
    "0910/0910.0127_arXiv.txt": {
        "abstract": "Optical/near-infrared observations for 14 globular cluster (GC) systems in early-type galaxies are presented. We investigate the recent claims (\\cite[Yoon, Yi \\& Lee 2006]{yoon06}) of colour bimodality in GC systems being an artefact of the non linear colour - metallicity transformation driven by the horizontal branch morphology. Taking the advantage of the fact that the combination of optical and near-infrared colours can in principle break the age/metallicity degeneracy we also analyse age distributions in these systems. ",
        "introduction": "The ubiquity of globular clusters (GCs) around all major galaxies and the fact that young star clusters are being observed in many mergers and starbursts we observe today suggests that GC formation traces important star formation events in the history of their host galaxies (\\cite[Brodie \\& Strader 2006]{bs06}).  GC systems generally exhibit a colour distribution that is bimodal (sometimes multimodal) in optical colour. Spectroscopic studies (eg. \\cite[Strader et al. 2005]{strader05}, \\cite[Cenarro et al. 2007]{cenarro07}) show that this bimodality is due to old subpopulations with ages greater than 10 Gyrs that differ in metallicity. Bimodal metallicity distributions may indicate distinct formation mechanisms, with age differences $\\sim2$\\,Gyrs still being allowed within the uncertainties on current age estimates. This metallicity bi/multimodality seems to hold for the majority of galaxies, except for the low mass galaxies which appear to have only a metal poor GC subpopulation. Recently \\cite[Yoon, Yi \\& Lee (2006)]{yoon06} challenged this metallicity bimodality interpretation saying that it is an artefact of the horizontal branch morphologies. These authors argue that non-linear colour metallicity relations may transform a unimodal metallicity distribution into a bimodal optical colour distribution. This issue has been investigated in more detail by \\cite [Cantiello \\& Blakeslee (2007)]{cb07} who concluded that combinations of optical and NIR colours are much less prone to this effect. Metallicity bimodality was thus contested and it is therefore important to investigate this topic in more detail. On the other hand, if metallicity bi/multimodality is genuine there could be multiple episodes/mechanisms of GC formation in a galaxy with a possible age difference. It is thus equally important to study the age distributions of these systems. It is still an open issue in the literature whether some early type galaxies known to contain old stellar populations from integrated light studies can host a high percentage of intermediate age (2-3 Gyrs) GCs. The two strongest cases are NGC\\,4365 (\\cite[Puzia et al. 2002]{puzia02}, \\cite[Larsen et al. 2005]{lbs05}) and NGC\\,5846 (\\cite[Hempel et al. 2003]{hempel03}). The comparison of 2-colour plots of the GC systems with simple stellar population (SSP) models (eg. \\cite[Bruzual \\& Charlot 2003]{bc03}, \\cite[Maraston 2005]{m05}) suggested the presence of intermediate age GCs in these galaxies.  With the aim of investigating ages and metallicity distributions in GC systems of early-type galaxies we examine 14 E/S0 systems, with $(m-M)<32$ and $M_B$  $<$ $-19$ through optical and NIR imaging. We obtained deep ($\\sim 3.4$ hours) K-band photometry with the LIRIS instrument in the William Herschel Telescope in La Palma which has a field of view of $4.2\\,\\times\\,4.2\\,arcmin^{2}$. These were combined with optical imaging extracted from the Hubble Space Telescope public archive in $g$ (F475W) and $z$ (F850LP) bands. GC candidates were selected automatically in galaxy subtracted images through standard routines. The following cuts were applied, defining the final sample: $g$ magnitude  $g\\,<\\,23$, optical colour $0.5\\,<\\,(g-z)\\,<\\,2$, effective radii $1\\,<\\,R_{eff}\\,<\\,15\\,pc$ and magnitude error $K_{err}\\,<\\,0.5$. With these restrictions, typical magnitude errors for the optical bands $g$ and $z$ came to be $\\sim\\,0.05$. For the NIR band the magnitude errors were measured to be $\\sim\\,0.06$ for a $K=19$ GC and $\\sim\\,0.2$ for a $K=20$ GC. Artificial star tests though show that the $K$-band magnitude errors can be underestimated by a factor of 2. In what follows it is shown what knowledge was gained from this data regarding the age distributions and the colour (metallicity) distributions of the GC systems studied here. ",
        "conclusions": "A sample of GC systems in E/S0s was observed in the K-band with LIRIS/WHT and combined with archival ACS/HST imaging in the $g$ and $z$ bands. From these observations we intended to study the overall ages and metallicity distributions of GC systems.  We find that age dating GCs as old or intermediate age depends on the choice of the model. Padova SSPs for old ages ($\\sim 12,14$ Gyrs) with \\cite[Marigo et al. (2008)]{marigo08} isochrones with a new treatment for the TP-AGB phase fits the data well. \\cite[Maraston (2005)]{m05} and \\cite[Charlot \\& Bruzual (2009)]{cb09} models yield intermediate ages  ($\\sim 2-3$ Gyrs) for the majority of the GCs. NGC\\,4365 GC system is not found to be an exception, in general all the GC systems in the 2-colour diagrams look very much alike. NGC\\,4406 is found to be the strongest candidate for having intermediate ages.  As far as metallicity distributions are concerned the bimodality becomes less evident in $(g-k)$ if compared to $(g-z)$ and even less pronounced in $(z-k)$. The latter colour is the one that should have the smallest contribution from horizontal branch stars. The disappearance of bimodality in this colour while evident in the optical $(g-z)$ colour could be attributed to the \\cite[Yoon, Yi \\& Lee (2006)]{yoon06} effect although the observational uncertainties could also account for it. The relation $(g-k)$ \\textit{vs.} $(z-k)$ shows a wiggle around  $(g-k)\\,\\sim\\,3.2$ and $(z-k)\\,\\sim\\,2$ which is present in the 14 Gyr SPoT-Teramo SSP models where there is a proper treatment of the horizontal branch morphology. It is concluded that redder colours such as $(z-k)$ are ideal metallicity indicators.  This study calls attention to the importance of choice of SSP models. How well a set of models fits the data seems to be linked to which stellar evolutionary stages are incorporated in it and how this is done. As far as deriving overall ages of GCs with the optical/NIR imaging technique, it becomes clear that different kinds of TP-AGB treatments yield different results. SSPs with recent TP-AGB models incorporated are consistent with old GC ages. The addition or not of a realistic treatment of the HB morphology can also make a big difference. Whether this evolutionary phase is or not responsible for optical colour bimodality cannot be answered with the data here presented."
    },
    "0910/0910.1317_arXiv.txt": {
        "abstract": "Analyses of the cosmic microwave background (CMB) radiation maps made by the Wilkinson Microwave Anisotropy Probe (WMAP) have revealed anomalies not predicted by the standard inflationary cosmology. In particular, the power of the quadrupole moment of the CMB fluctuations is remarkably low, and the quadrupole and octopole moments are aligned mutually and with the geometry of the Solar system. It has been suggested in the literature that microwave sky pollution by an unidentified dust cloud in the vicinity of the Solar system may be the cause for these anomalies. In this paper, we simulate the thermal emission by clouds of spherical homogeneous particles of several materials. Spectral constraints from the WMAP multi-wavelength data and earlier infrared observations on the hypothetical dust cloud are used to determine the dust cloud's physical characteristics. In order for its emissivity to demonstrate a flat, CMB-like wavelength dependence over the WMAP wavelengths (3 through 14~mm), and to be invisible in the infrared light, its particles must be macroscopic. Silicate spheres from several millimetres in size and carbonaceous particles an order of magnitude smaller will suffice. According to our estimates of the abundance of such particles in the Zodiacal cloud and trans-neptunian belt, yielding the optical depths of the order of $10^{-7}$ for each cloud, the Solar-system dust can well contribute $10\\;\\mu$K (within an order of magnitude) in the microwaves. This is not only intriguingly close to the magnitude of the anomalies (about $30\\;\\mu$K), but also alarmingly above the presently believed magnitude of systematic biases of the WMAP results (below 5~$\\mu$K) and, to an even greater degree, of the future missions with higher sensitivities, e.g.\\ PLANCK. ",
        "introduction": "The five-year release~\\citep{Hinshaw-et-al-2009} of the Wilkinson Microwave Anisotropy Probe \\citep[WMAP,][]{Bennett-et-al-2003ApJ} data represents the current highlight of a project which provided a so far unprecedented amount of high precision cosmological data. On small angular scales the measurement of the cosmic microwave background (CMB) anisotropies led to the precise confirmation of the cold dark matter ($\\Lambda$CDM) model of a nearly spatially flat universe, dominated by dark energy and non-baryonic dark matter \\citep{Spergel-et-al-2006astro.ph}. Nevertheless at large scales a number of unexpected, potentially damaging anomalous features in the CMB have been reported, indicating that either our current understanding of standard inflationary cosmology or the data processing, including foreground removal techniques are as yet inadequate.  Among these anomalies, the expansion of the CMB sky in spherical harmonic functions resulted in an octopole which is unusually planar and oriented parallel to the quadrupole \\citep{Tegmark-et-al-2003PhRvD,deOliveiraCosta-et-al-2004PhRvD}. Three of the four planes determined by the quadrupole and octopole are orthogonal to the ecliptic at 99.1\\% confidence level (CL), and the normals to these planes are aligned with the direction of the cosmological dipole and with the equinoxes, inconsistent with Gaussian random, statistically isotropic skies at 99.8\\% CL \\citep{Schwarz-et-al-2004PhRvL}.  The presence of non-Gaussian features in the CMB temperature fluctuations was reported by \\citet{Copi-et-al-2004}. \\citet{Land-Magueijo-2005} suggested that the presence of preferred directions in the low-order multipoles also extends to higher multipoles beyond the octopole.  Of course any of these anomalies challenges the validity of the standard scenario of inflationary cosmology which predicts scale-free, statistically isotropic and Gaussian random CMB temperature fluctuations and uncorrelated multipoles. However, although it is very unlikely that these features are just a statistical fluke, their cosmological origin is still an open debate. Therefore, besides various kinds of potential new physics \\citep{Hannestad-MersiniHoughton-2005,Moffat-2005,Jaffe-et-al-2005,Gordon-et-al-2005}, a compact cosmic topology \\citep{deOliveiraCosta-et-al-2004PhRvD,Mota-et-al-2004,Cornish-et-al-2004} or modified inflation \\citep{Piao-2005,Linde-2004,Hunt-Sarkar-2004} and also conventional effects of Galactic foreground emission \\citep{Eriksen-et-al-2004,Slosar-Seljak-2004,Naselsky-et-al-2005} have been suggested as possible physical explanations.  The significance of the correlations of the quadrupole and octopole with the ecliptic plane, however, hints at the Solar system as the origin of an unaccounted bias~\\citep{Schwarz-et-al-2004PhRvL,Copi-et-al-2006}. \\cite{Starkman-Schwarz-2005} speculated that an unknown dark cloud of dust in the solar neighbourhood may have contributed to the microwave sky surveyed by the WMAP. \\cite{Frisch-2005} proposed that the interstellar dust trapped magnetically in the heliosphere can possibly explain the CMB anomalies. \\cite{Babich-et-al-2007ApJ} discussed the possibility of discovering the dust of the trans-neptunian belt in the WMAP data. None of the investigations, however, gave fully convincing explanations to solve the problem of the CMB anomalies.  This paper opens a series of publications in which we simulate the microwave thermal emission of dust clouds inside or in the vicinity of the Solar system, and test whether these clouds can be responsible for the unexplained correlations {or can be discovered in the CMB experiment data}. Here we focus on the spectral constraints on material and size distribution of the particles forming a cloud invisible in the infrared wavelengths and imitating closely the CMB emission in the microwaves. Such cloud would indeed contribute to the CMB maps without us knowing about it. The absolute photometry and spatial geometry of candidate clouds' contributions to the CMB maps are left for further scrutiny in follow-up studies.  {The paper is organised as follows. Section~\\ref{excess temp} deals with the general constraints on an unknown cloud's spectrum, from the infrared wavelengths to microwaves. It shows how a cloud of macroscopic particles can be noticeable in the microwaves and still avoid detection in the infrared light. Candidate clouds are proposed in the Solar system that can add non-negligible emission in the microwaves without being easily recognized in the infrared light. Estimates of the microwave temperatures of some known dust clouds are made, supporting the case for a detailed study. Constraints on the microwave spectra of foreground sources from the Internal Linear Combination (ILC) maps derived from the WMAP multi-wavelength data \\citep{Bennett-et-al-2003ApJS-maps} are also formulated to facilitate the determination of plausible composition and size distribution of the cloud. Section~\\ref{mie scattering} introduces the elements of the Mie theory of light scattering necessary for our calculations of the thermal emission from dust particles, and borrows bibliographic sources on the optical properties of several chemical compositions from the database \\citep{Henning-et-al-1999}. Thermal emission spectra of various sample and one natural dust clouds are generated in Sec.~\\ref{temp spectra} and tested against the constraints placed by the ILC maps. Conclusions are drawn in Sec.~\\ref{conclusions}.} ",
        "conclusions": "We have investigated the microwave thermal emission by dust in the Solar system. We applied the Mie theory of light scattering by spherical homogeneous particles in order to characterise the thermal emission spectra of silicate, carbonaceous, iron and icy particles. Our study is partly motivated by the WMAP observations of the CMB fluctuations that revealed large-scale structures aligned with the Solar-system geometry which are difficult to explain by the standard inflationary cosmology. One possibility to produce such fluctuations is by a dust cloud inside or in the vicinity of the Solar system. Another motivation for this study is to assess the feasibility of detection of dust in and near the Solar system in the microwaves.  We used the WMAP multi-wavelength observations and infrared surveys to constrain the physical properties of particles constituting the hypothetical cloud and to estimate the microwave emission by solar-system dust. We have found that only macroscopic, mm-sized silicate or carbonaceous grains could produce thermal emission with a spectrum compatible with that of the CMB fluctuations. Smaller dust grains, as well as iron particles, emit with a spectrum that can easily be distinguished from the CMB. The small particles would also be so bright in the infrared light that they would have been seen by the relevant telescopes, whereas the big particles have a flat emissivity throughout the spectrum, so that the abundant small grains outshine them at infrared wavelengths.  In order to attain the flat emissivity at the WMAP wavelengths, silicate particles must be several mm in size at least, whereas carbonaceous particles can be an order of magnitude smaller. This makes the carbonaceous particles the most likely candidate, as small grains are typically more likely in dust clouds than big ones. Meteoroids from short-period comets~\\citep{Hughes-McBride-1990} are a plausible candidate for such cloud. The dust needs not necessarily be cold nor remote. When the cloud is composed of particles with a broad size distribution, even smaller dust grains can contribute to an overall flat spectrum without revealing themselves. For example, all particles of the \\cite{Gruen-et-al-1985} meteoroid flux model with sizes above 100~$\\mu$m have an integral spectrum sufficiently flat to pass the ILC procedure used to clean the WMAP observations from biases, if they are composed of carbonaceous material. The trans-neptunian belt is another plausible candidate. According to our preliminary estimates, each candidate cloud can emit roughly $\\sim10~\\mu$K in the microwaves.  We found studies \\citep[see][for references]{Stognienko-et-al-1995} showing that the Mie theory provides in general more pessimistic estimates of compatibility between the dust and CMB spectra than advanced light scattering theories taking nonsphericity and inhomogeneity into account. It means that the real particles can be even smaller than those quoted above in this conclusion.  There is a lack of measurements of the optical constants of carbonaceous particles above a wavelength of 2~mm. Given that this material provides the best chances so far to explain the WMAP anomaly, laboratory measurements of carbonaceous particles would be very helpful for further studies of the problem.  In subsequent papers of this series, using the thermal emission model introduced here, we will test the absolute photometry and projected geometry of various known and hypothetical dust clouds in and near the Solar system, against the maps of CMB fluctuations made by WMAP."
    },
    "0910/0910.4018_arXiv.txt": {
        "abstract": "We analyze the projected axial ratio distribution, $p(b/a)$, of galaxies that were spectroscopically selected from the Sloan Digital Sky Survey (DR6) to have low star formation rates.  For these quiescent galaxies we find a rather abrupt change in $p(b/a)$ at a stellar mass of $\\sim10^{11}\\msol$: at higher masses there are hardly any galaxies with $b/a<0.6$, implying that essentially none of them have disk-like intrinsic shapes and must be spheroidal. This transition mass is $\\sim 3-4$ times higher than the threshold mass above which quiescent galaxies dominate in number over star-forming galaxies, which suggests that these mass scales are unrelated.  At masses lower than $\\sim 10^{11}\\msol$, quiescent galaxies show a large range in axial ratios, implying a mix of bulge- and disk-dominated galaxies.  Our result strongly suggests that major merging is the most important, and perhaps only relevant, evolutionary channel to produce massive ($>10^{11}\\msol$), quiescent galaxies, as it inevitably results in spheroids. ",
        "introduction": "Even galaxies with little star formation activity continue to evolve, as evidenced by the substantial increase of their cosmic stellar mass density over the past 7 billion years \\citep{bell04b, faber07, brown07}.  This must be related to the decreasing star formation activity over the same period \\citep[e.g.,][]{lefloch05}, and the production of such quiescent galaxies through the truncation of star formation \\citep[e.g.,][]{faber07, bell07}; the color scatter among quiescent galaxies and its evolution are in precise agreement with such a scenario \\citep{ruhland09}.  There are, however, quiescent galaxies at all redshifts $z\\lesssim 1.3$ that are more massive than the most massive star-forming galaxies. This implies that star formation in the most massive galaxies was truncated even earlier, and/or that mergers play an important role in producing massive galaxies.  Evidence for the early formation of massive galaxies is provided by their old stellar populations. However, we need to bear in mind that there can be a large difference between the age of the stellar population and the assembly age, especially if mergers are important, as is the case in a hierarchical framework for galaxy formation \\citep{delucia07}.  Hence, the number density evolution of galaxies is important in constraining their assembly history.  Measuring this is difficult because of its sensitivity to the luminosity evolution correction, especially for massive galaxies at the exponential cut-off of the mass function.  As a result, there is no consensus among the currently available measurements \\citep{cimatti06, wake06, brown07, cool08}.  Given these difficulties, other observations have been used to either directly or indirectly constrain the assembly of galaxies.  Merging activity among the massive galaxy population is observed \\citep[e.g.,][]{vandokkum99, vandokkum05, bell06a, bell06b, lin08}, and has been shown to produce a color-magnitude relation that is in agreement with observations \\citep{skelton09}.  However, its cosmological relevance has always been difficult to determine, given the uncertainties in converting observed merger fractions to merger rates and the associated growth in mass.  An independent and indirect indication that massive galaxies undergo continuous evolution is provided by the recent result that high-redshift quiescent galaxies are substantially smaller than local galaxies with the same mass \\citep[see,][and references therein]{vanderwel08c}.  This strongly suggests that mergers are important \\citep[see, e.g.,][]{vanderwel09a}, and that the assembly of massive galaxies is continuing up until the present day.  Another indirect, yet powerful, constraint is provided by the evolution in the clustering and halo occupation distribution of red galaxies \\citep{white07, conroy07, brown08}: the evolution in the clustering strength of red galaxies is slower than expected in the absence of merging.  In this Letter we address the question whether major merging is the dominant mechanism for the production of very massive, quiescent galaxies.  The argument that we invoke is simply that major merging generally leads to rounder galaxies.  An analysis of the shape distribution of quiescent galaxies can therefore constrain the importance of merging.  Since merging among galaxies with mass ratios of $\\lesssim 3$ is the only known mechanism to produce round galaxies (see Section 3 for further discussion), this is a powerful test.  The disadvantage of this method, compared to those mentioned above, is that no information about the time scale and epoch of galaxy assembly can be inferred.  \\citet{vincent05} and \\citet{padilla08} were the first to systematically study the axial ratio distribution, $p(b/a)$, of a large number of galaxies, selected from the Sloan Digital Sky Survey (SDSS).  Through a detailed analysis, they infer the intrinsic shape distribution and the effect of extinction.  Both divide the sample into 'elliptical' and 'spiral' galaxies, and confirmed that luminous 'elliptical' galaxies are, on average, rounder and tri-axial, compared to low-luminosity 'ellipticals', which are more elongated and oblate \\citep{davies83, franx91}, and display disky isophotes \\citep{jorgensen94}. This phenomenon is not recent: \\citet{holden09a} showed that this trend persists at least out to $z\\sim 1$.  Here we present a complementary, modified analysis, focusing on $p(b/a)$ as a function of stellar mass for quiescent, i.e., non-star-forming, galaxies.  Because mass-to-light ratios are well constrained by broad-band colors for quiescent galaxies, stellar mass estimates are robust.  This is essential for our purposes, as we are interested in the most massive objects, i.e., those that populate the exponential tail of the mass function.  Furthermore, as opposed to previous studies, we pre-select galaxies independent of their photometric properties.  Our shape-independent, spectroscopic selection criteria circumvent the biases that are potentially introduced by selecting galaxies by their 'morphological' properties, or some pre-defined surface brightness profile.  With this sample, for which we have determined axial ratios from our own fits to two-dimensional light distributions, we address the following specific questions. Are high-mass, quiescent galaxies rounder than low-mass quiescent galaxies? If so, is there a mass limit at which $p(b/a)$ distinctly changes, and above which disk-dominated are completely absent?  Such evidence would imply that the only evolutionary path to such masses is a disk-destroying mechanism, i.e., major merging. ",
        "conclusions": "\\label{res}  In Figure \\ref{M_q}(a) we show $p(b/a)$ of the 17,480 spectroscopically selected, quiescent galaxies as a function of stellar mass.  $p(b/a)$ is shown in gray scale, with the percentiles of the cumulative $b/a$ distribution shown as (red) lines.  Figure \\ref{M_q}(a) immediately demonstrates that for quiescent galaxies, the projected axial ratio distribution is a strong function of stellar mass. In the narrow mass range $8\\times 10^{10} \\lesssim M_*/M_{\\odot} \\lesssim 2\\times 10^{11}$ there is a rapid decrease in the number of galaxies with small axial ratios.  As further illustrated by Figure \\ref{M_q}(b), above $M_*\\sim 2\\times 10^{11}~M_{\\odot}$ quiescent galaxies with $b/a<0.6$ are essentially absent.  This result shows that evolutionary paths that lead to quiescent galaxies with stellar mass $M_* \\gtrsim 2\\times 10^{11}~M_{\\odot}$ all but exclude the existence, or the survival, of highly flattened, disk-like stellar components.  As highly flattened stellar systems are quite common at lower masses, in the possible realm of plausible progenitors of high-mass galaxies, this result implies the destruction of the flattened component in whatever process causes growth beyond $M_*\\sim 2\\times 10^{11}~\\msol$.  Therefore, our result that essentially all quiescent galaxies with masses larger than $M_*\\sim 2\\times 10^{11}~\\msol$ are round strongly suggests that for such galaxies major mergers are the dominant, perhaps even unique, formation channel.  The destruction of a stellar disk requires a major merger, i.e., a merger involving progenitors with a relatively small mass ratio of at most $\\sim 3$, mergers with a larger mass ratio leaving stellar disks intact \\citep[see, e.g.,][]{bekki98, bournaud04}.  Moreover, most likely, the progenitors are not very gas rich, as this would produce a disky remnant \\citep[e.g.,][]{naab06b}.  It has been suggested that cold flows are responsible for the formation of massive, classical bulges at high redshift \\citep{dekel09, ceverino09}.  In this scenario, intensely star-forming 'knots' merge, forming a massive bulge \\citep[see also][]{noguchi99}. However, a substantial fraction of the mass ($\\sim 50\\%$) is still predicted to reside in a disk.  Even at later stages, when the gas disk has become stable against fragmentation and collapse \\citep[``morphological quenching'';][]{martig09}, the stellar disk remains intact and contains a non-negligible fraction of the total mass.  In short, although cold flows plausibly produce quiescent galaxies, the end-products will not be uniquely round.  Only in the case of sufficient merger activity would galaxies become spheroidal.  In passing, we note that the sharp decrease in the fraction of very round galaxies ($b/a > 0.8$) at the very highest masses ($M>3\\times 10^{11}~\\msol$) signifies that such high mass galaxies are typically brightest group/cluster galaxies, which tend to be slightly more elongated than 'normal' massive elliptical galaxies \\citep[see,][]{bernardi08}.  As already noted above, at masses lower than $M_*\\sim 10^{11}~\\msol$, quiescent galaxies display a large range in axial ratios, which implies that star formation truncation mechanisms below $10^{11}~\\msol$ are often not associated with the destruction of the disk.  It remains to be tested whether $p(b/a)$ of low-mass quiescent galaxies is similar to or different from $p(b/a)$ of star-forming galaxies in the same mass range. Such an analysis, which is non-trivial because of the effects of extinction and color gradients, will constrain the degree to which mergers or bulge growth regulate star formation at these lower masses.  Another open question concerns the number and properties of massive, star-forming galaxies.  Morphological studies suggest that a large fraction ($20\\%-40\\%$) of all galaxies more massive than $M\\sim 10^{11}~M_{\\odot}$ are late-type galaxies \\citep{vanderwel08a, bamford09}, and their high masses of at least some of these objects are confirmed by their rotational velocities \\citep[e.g.,][]{courteau07}. Yet, the degree to which such galaxies are disk-dominated and should be considered actively star forming remains to be determined.  It would, therefore, be premature to conclude that merging is the only way to produce a massive galaxy in general, and therefore restrict this proposition to the formation of massive, quiescent galaxies.  The picture sketched by the axial ratio distribution of quiescent galaxies is in agreement with the strong correlation between structure and mass for galaxies in general \\citep[e.g.,][]{kauffmann03b, vanderwel08a} and early-type galaxies in particular \\citep[e.g.,][]{caon93, graham03}: high-mass galaxies are more concentrated and have higher S\\'ersic indices than low-mass galaxies. These trends are an indirect indication of a decreasing importance of disks for galaxies with higher masses, although part of this trend is caused by the increase in S\\'ersic index with galaxy mass among spheroidal galaxies.  In our sample we see a similar trend: in the mass range $10^{10} < M_*/\\msol < 2\\times 10^{10}$, 41\\% of the galaxies have S\\'ersic indices $n<3$, whereas at higher masses, $M_*>2\\times 10^{11}\\msol$, only 3\\% have such low S\\'ersic indices. We postpone a full exploration of the joint behavior of shape and structure as a function of galaxy mass until a future paper, but it is encouraging that the apparent absence of prominent disks in high-mass, quiescent galaxies is reflected in both the S\\'ersic index and $p(b/a)$."
    },
    "0910/0910.4532_arXiv.txt": {
        "abstract": "We consider the early universe at temperatures close to the fundamental scale of gravity ($M_D\\ll M_{Planck}$) in models with extra dimensions. At such temperatures a small fraction of particles will experience transplanckian collisions that may result in microscopic black holes (BHs). BHs colder than the environment will gain mass, and as they grow their temperature drops further. We study the dynamics of a system (a {\\it black hole gas}) defined by radiation at a given temperature coupled to a distribution of BHs of different mass. Our analysis includes the production of BHs in photon-photon collisions, BH evaporation, the absorption of radiation, collisions of two BHs to give a larger one, and the effects of the expansion. We show that the system may follow two different generic paths depending on the initial temperature of the plasma. ",
        "introduction": "Large (ADD) \\cite{ArkaniHamed:1998rs} or warped (RS) \\cite{Randall:1999ee} extra dimensions open the possibility that the fundamental scale of gravity $M_D$ is much lower than $M_P\\approx 10^{19}$ GeV. This could imply that the {\\it transplanckian} regime \\cite{Banks:1999gd} is at accesible energies. Collisions in that regime are very different from what we have experienced so far in particle colliders: due to its spin 2 the graviton becomes strongly coupled and dominates at distances that increase with the center of mass energy $\\sqrt{s}$. In particular, at small impact parameters one expects that gravity {\\it bounds} the system and the two particles collapse into a microscopic black hole (BH) of mass around $\\sqrt{s}$. Such BH would evaporate \\cite{Hawking:1974rv} very fast (the time scale is set by $1/M_D$) into final states of high multiplicity, a possible LHC signature that has been extensively discussed in the literature \\cite{Dimopoulos:2001hw}.  These collisions, however, may have also occurred in the early universe if the temperature was ever close to $M_D$. Consider a period of inflation produced by a field on the brane, followed by reheating into 4--dimensional (4-dim) species. If the reheating temperature is $T_{rh}\\ge 0.1 M_D$, particles in the tail of the Boltzmann distribution may collide with enough energy to form BHs. Hot BHs will evaporate, but if a mini-BH is colder than the environment it absorbs more than it emits \\cite{Majumdar:2002mra}, growing and becoming colder, which in turn increases its absorption rate (see below). The fact that heavier BHs live longer distinguishes them from massive particles or string excitations possibly produced at these temperatures, since the lifetime of the latter is inversely proportional to their mass. As a consequence, one expects that BHs are a critical ingredient at temperatures near the fundamental scale. Notice that these BHs are not the primordial ones formed from the gravitational collapse of density fluctuations \\cite{Carr:1974nx,Barrow:1991dn,Argyres:1998qn,Majumdar:2005ba} at lower temperature $T$.  Although BHs would also be present in a 4-dim universe at $T\\sim M_P$, they are most relevant in TeV gravity models. The reason is that the expansion of the universe is a long-distance process, so its rate is dictated by $M_P$ (not $M_D$) also in models of TeV gravity. If the {\\it bulk} is basically empty \\cite{ArkaniHamed:1998nn,Hannestad:2001nq,Starkman:2000dy} or {\\it thin} (see below), then the expansion at $T\\le M_D$ will be negligible in terms of the fundamental time scale $1/M_D$, and there is plenty of time for collisions producing BHs to occur and for the BHs to grow. In contrast, in a 4dim universe the expansion rate at $T\\approx M_P$ is of order $H^{-1}\\approx 1/M_P$, the temperature of the universe drops before BHs have grown, and once it goes below $T_{BH}$ they evaporate in the same time scale.  In this article we explore the implications of having initial temperatures near the fundamental scale of gravity. First we define a consistent set up for TeV gravity. Then we study the behaviour of a single BH inside a thermal bath in an expanding universe. Finally we consider the generic case, a {\\it black hole gas}, with radiation coupled to a distribution of BHs of different mass. We find remarkable that the effect of these mini-BHs in the early universe has been almost completely overlooked in the literature (the only analysis that we have found is given by Conley and Wizansky in \\cite{Conley:2006jg}), although many authors have considered temperatures close (and above) the Planck scale (see \\cite{Brandenberger:1988aj} and references therein). ",
        "conclusions": "A transplanckian regime at accessible energies would have new and {\\it peculiar} implications in collider physics and cosmology. What makes this regime special is that as the energy grows, softer physics dominate. For example, the production of a {\\it regular} heavy particle at the LHC would provide an event with very energetic (hard) jets from its decay. In contrast, the production of a mini-BH would be seen as a high multiplicity event, with dozens of less energetic (soft) jets. Analogously, the production of massive particles or mini-BHs in the early universe would have very different consequences. The heavier the elementary particle, the shorter its lifetime, which tends to decouple these particles from low temperatures. For mini-BHs is just the opposite, heavier BHs are colder and live longer. In addition, if an approximate symmetry makes a massive relic long lived, its late decay will produce very energetic particles. BHs would imply a much softer spectrum of secondaries \\cite{Draggiotis:2008jz}, with different cosmological consequences.  In this paper we have analyzed the dynamics of a two-component gas with photons and BHs in an expanding universe. The system is characterized by the temperature $T(t)$ of the radiation and the distribution $f(M,t)$ of BHs. Our equations take into account BH production, absorption of photons, BH evaporation, and collisions of two BHs. We have discussed the two generic scenarios that may result from an initial temperature close to the fundamental scale of gravity $M_D$. In the first scenario BHs {\\it empty} of radiation the universe and dominate $\\rho$ before a Hubble time, whereas in the second case there are few BHs that grow at constant $T$ up to $t\\approx H^{-1}$.  We think that the work presented here is a necessary first step in the search for observable effects from these BHs. It could also lead us to a better understanding of the early universe at the highest temperatures."
    },
    "0910/0910.3859_arXiv.txt": {
        "abstract": "{During solar flares a large number of charged particles are accelerated to high energies, but the exact mechanism responsible for this is, so far, still unclear. Acceleration in collapsing magnetic traps is one of the mechanisms proposed.  }{In the present paper we want to extend previous 2D models for collapsing magnetic traps to 3D models and to 2D models with shear flow.}{We use analytic solutions of the kinematic magnetohydrodynamic (MHD) equations to construct the models. Particle orbits are calculated using the guiding centre approximation.}{We present a general theoretical framework for constructing kinematic MHD models of collapsing magnetic traps in 3D and in 2D with shear flow. A few illustrative examples of collapsing trap models are presented, together with some preliminary studies of particle orbits. For these example orbits, the energy increases roughly by a factor of 5 or 6, which is consistent with the energy increase found in previous 2D models.}{} ",
        "introduction": "One of the main features of solar flares is the acceleration to high energies of a substantial number of charged particles within a short period of time. The explanation of how this happens is one of the most important open questions in solar physics. There is general agreement that the energy released in solar flares is previously stored in the magnetic field, but the exact physical mechanisms by which this energy is released and converted into bulk flow energy, thermal energy, non-thermal energy and radiation energy are still a matter of discussion \\citep[e.g.][]{miller:etal97,aschwanden02,neukirch05b,neukirch:etal07,krucker:etal08,aschwanden09}. Using observations of non-thermal high-energy (hard X-ray and $\\gamma$-ray) radiation, it is estimated that a large fraction of the released magnetic energy (up to  the order of 50 \\%)  is converted into non-thermal energy in the form of high energy particles \\citep[e.g.][]{emslie:etal04,emslie:etal05}.  A variety of possible particle acceleration mechanisms have been suggested including direct acceleration in the parallel electric field associated with the reconnection process, stochastic acceleration by turbulence and/or wave-particle resonance, shock acceleration or acceleration in the inductive electric field of the reconfiguring magnetic field \\citep[see e.g.][for a detailed discussion and further references]{miller:etal97,aschwanden02,neukirch05b,neukirch:etal07,krucker:etal08}. So far, none of the proposed mechanisms can explain the high-energy particle fluxes within the framework of the standard solar flare thick target model. This has recently prompted suggestions of alternative acceleration scenarios \\citep[e.g.][]{fletcher:hudson08,birn:etal09}.  \\citet{somov92} and \\citet{somov:kosugi97} suggested that the reconfiguration of the magnetic field during a flare could contribute to the acceleration of particles. Due to the geometry of the magnetic field charged particles could be trapped while the magnetic field lines relax dynamically. In such a collapsing magnetic trap (CMT from now on) the kinetic energy of the particles could increase due to the betatron effect, as the magnetic field strength in the CMT increases, and due to first-order Fermi acceleration, as the distance between the mirror points of particle orbits decreases due to the shortening of the field lines. There is also some observational evidence of post-flare field lines relaxation (field line shrinkage) from Yohkoh \\citep[e.g.][]{forbes:acton96} and Hinode \\citep[e.g.][]{reeves:etal08a} observations.  Various fundamental properties of the particle acceleration process in CMTs have been investigated by Somov and co-workers \\citep[e.g.][]{bogachev:somov01,bogachev:somov05,bogachev:somov09,kovalev:somov02,kovalev:somov03a,kovalev:somov03b,somov:bogachev03}, including the relative efficiencies of betatron and Fermi acceleration,  the effect of collisions, the role of velocity anisotropies  and the evolution of the energy distribution function in a CMT. In all cases a basic model for CMTs has been used. \\citet{karlicky:kosugi04} also investigated particle acceleration, plasma heating and the resulting X-ray emission using a simple CMT model and a simplified equation of motion for the particles. \\citet{karlicky:barta06} used CMT-like electromagnetic fields taken from an MHD simulation of a reconnecting current sheet to investigate acceleration using test particle calculations with a view to explain hard X-ray loop-top sources.  Some indication that CMTs might be relevant for X-ray loop top sources has been provided by \\citet{veronig06}. A very simple time-dependent trap model was also used by \\citet{aschwanden04} to explain the pulsed time profile of energetic particle injection during flares.  A general theoretical framework for more detailed analytical CMT models based on kinematic MHD, i.e. with given bulk flow profile, in Cartesian coordinates was presented by \\citet{giuliani:etal05}  for 2D and 2.5D magnetic fields, but excluding flow in the invariant direction. Some examples of model CMTs were given together with a calculation of a particle orbit based on non-relativistic guiding centre theory \\citep[see e.g.][]{northrop63}. It was found that in the models studied the curvature drift and the gradient-$B$ drift play an important role in the acceleration process. Similar findings have also been made, albeit in systems of a much smaller length, in the investigation of particle acceleration in particle-in-cell simulations of collisionless magnetic reconnection \\citep[e.g.][]{hoshino:etal01}.  The advantage of kinematic MHD models compared to e.g. MHD simulations is that they allow us to obtain analytical expressions for the electromagnetic fields of the CMT. This makes the integration of particle orbits more accurate, because there is no need for interpolation of the fields between grid points. Furthermore the investigation of different model features is possible in an easy way by varying model parameters. The major disadvantage of kinematic MHD models is their lack of self-consistency, but this is not too critical for the purpose of test particle calculations.  Particle acceleration through rapid reconfiguration of the magnetic field has also been identified as one of the mechanisms for particle energization during magnetospheric substorms \\citep[e.g.][]{birn:etal97,birn:etal98,birn:etal04}. During a substorm the stretched magnetic field of the magnetotail reconnects, leading to a so-called dipolarisation of the near-Earth tail, which is in principle very similar to the evolution of the magnetic field in a CMT associated with a solar flare. A general comparison of flare and substorm/magnetotail phenomena based on observations has recently been presented by \\citet{reeves:etal08b}.  The purpose of the present paper is to extend the theoretical framework for kinematic MHD CMT models given by \\citet{giuliani:etal05} to 2.5D models with flow in the invariant direction and to fully three-dimensional models. This is necessary for a number of reasons:  \\begin{enumerate}  \\item The theory of kinematic MHD CMTs as developed so far by \\citet{giuliani:etal05} only allows for a magnetic field component in the invariant direction, but not for a component of the flow velocity in this direction. Without this component of the flow velocity the magnetic field component in the invariant direction can only increase in a CMT due to magnetic flux conservation. It is, however, to be expected that during a flare magnetic shear will be reduced rather than increased and therefore the introduction of a component of the flow velocity in the invariant direction is a necessary extension to be able to make the 2.5D models more realistic.  \\item In the 2D cases investigated by \\citet{giuliani:etal05} the acceleration due to curvature and gradient-B drift occurs in the invariant direction. This is due to the fact that the particles gain energy while moving parallel or anti-parallel (in the case of electrons) to the inductive electric field which in a 2D trap is in the invariant direction. Due to the spatial symmetry the electric field does not vary in this direction and this will have an influence on the acceleration process. It is therefore important to investigate the differences of the acceleration process between 2D models and non-symmetric 3D CMT models in the future.  \\item \\citet{giuliani:etal05} have already discussed a possible way of extending the 2D theory to three dimensions using Euler potentials. While Euler potentials allow a relatively straightforward extension of the  theory to 3D by simple analogy to the 2D case, they are not easy to use in the modelling process, which is already intrinsically more difficult in three dimensions. We therefore present in this paper an extension to the theory which makes it possible to avoid the explicit calculation of Euler potentials and uses the magnetic field directly.  \\end{enumerate}  The paper is organised as follows. In Sect. \\ref{sec:2d_theory} we briefly summarise the present state of the kinematic MHD theory of CMTs, before presenting its extensions to 2.5D with shear flow and to 3D. In Sect. \\ref{sec:examples} a couple of  illustrative examples of CMT models based on the new theoretical descriptions are shown, followed by examples of test particle calculations in Sect. \\ref{sec:orbits}. We conclude the paper in Sect. \\ref{sec:summary} with a summary and conclusions. Appendix A gives more detail of the calculation of the 3D field using Euler Potentials. ",
        "conclusions": "\\label{sec:summary}  We have developed a fully analytical model for kinematic time-dependent, 2D and 3D collapsing magnetic traps. This kinematic approach has the advantage that it allows us full control over all the features of the model, but has the disadvantage the modelling of the plasma system is not self-consistent.  In the present paper, we have shown how to build kinematic CMT models using the magnetic field directly, rather than using a flux function or Euler potentials. This is much easier and more straightforward to use, especially in 3D, than the theory presented in \\citet{giuliani:etal05}.  We have given illustrative examples of collapsing traps with transformations that give rise a shear flow in 2D and magnetic twist in 3D. We have calculated particle orbits for these new CMT models using guiding centre theory. For those orbits, the CMT models were found to give similar relative energy gains as with the 2D CMT without shearing. The particle orbits are different from the 2D CMT model by \\citet{giuliani:etal05} despite starting from the same initial position due to the differences in field line motion caused by the shear flow in 2D and by the twisting motion in 3D. The examples shown in this paper have been chosen specifically to be comparable with the example shown in \\citet{giuliani:etal05}. Different CMT models could allow for higher energy gains and  will be considered in future work. There are also many other possible combinations of initial positions, initial particle energy and pitch angles, as well as investigating proton/ion orbits as well as electron orbits. A systematic investigation for the 2D model of \\citet{giuliani:etal05} has shown that energy gain factors of order 50 or higher are possible for that model \\citep{grady:etal09}. A similar investigation is planned for the future for 2D with shear flow CMTs and 3D CMTs using the theory presented in this paper."
    },
    "0910/0910.1547_arXiv.txt": {
        "abstract": "In this paper we examine the effect of X-ray and Ly$\\alpha$ photons on the intergalactic medium temperature. We calculate the photon production from a population of stars and micro-quasars in a set of cosmological hydrodynamic simulations which self-consistently follow the dark matter dynamics, radiative processes as well as star formation, black hole growth and associated feedback processes.  We find that, {\\it (i)} IGM heating is always dominated by X-rays unless the Ly$\\alpha$ photon contribution from stars in objects with mass $M<10^8$~M$_\\odot$ becomes significantly enhanced with respect to the X-ray contribution from BHs in the same halo (which we do not directly model). {\\it (ii)} Without overproducing the unresolved X-ray background, the gas temperature becomes larger than the CMB temperature, and thus an associated 21~cm signal should be expected in emission, at $z \\simlt 11.5$. We discuss how in such a scenario the transition redshift between a 21~cm signal in absorption and in emission could be used to constraint BHs accretion and associated feedback processes. ",
        "introduction": "While the investigation of the very high-$z$ universe, $z \\sim 1000$, is possible through the Cosmic Microwave Background (CMB) radiation and detection of a handful of objects with redshift as high as $\\sim 7$ has recently been possible, there is a lack of observational data in the redshift interval $7 \\simlt z \\simlt 1000$. To detect radiation from this interval, telescopes with exceptional sensitivity in the IR and radio bands are needed. {\\tt JWST}\\footnote{http://ngst.gsfc.nasa.gov} (James Webb Space Telescope), for example, with its nJy sensitivity in the $1-10$~$\\mu$m infrared regime, is ideally suited for probing optical-UV emission from sources at $z>10$. Similarly, the planned generation of radio telescopes as {\\tt SKA}\\footnote{http://www.skatelescope.org} (Square Kilometer Array), {\\tt LOFAR}\\footnote{http://www.lofar.org} (LOw Frequency ARray), {\\tt 21cmA} (21cm Array) and {\\tt MWA}\\footnote{http://web.haystack.mit.edu/arrays/MWA/} (Murchison Widefield Array) will open a new observational window on the high redshift universe. In particular, the detection of the 21~cm line associated with the hyperfine transition of the ground state of neutral hydrogen, holds the promise to shed light on the reionization process and its sources.  This observational progress has been accompanied by a flourishing of theoretical activity, aimed at providing predictions for the planned generation of telescopes.  It has long been known (e.g. \\citealt{Field_59}) that neutral hydrogen in the intergalactic medium (IGM) and gravitationally collapsed systems may be directly detectable in emission or absorption against the CMB at the frequency corresponding to the redshifted HI 21~cm line (associated with the spin-flip transition from the triplet to the singlet ground state). \\citet{MadauMeiksinRees_97} first showed that 21~cm tomography could provide a direct probe of the era of cosmological reionization and reheating. In general, 21~cm spectral features will display angular structure as well as structure in redshift space due to inhomogeneities in the gas density field, hydrogen ionized fraction, and spin temperature. Several different signatures have been investigated in the recent literature among which fluctuations in the 21~cm line emission induced by the ``cosmic web'' \\citep{Tozzi_etal_00}, by the neutral hydrogen surviving reionization (e.g. \\citealt{CiardiMadau_03, FurlanettoSokasianHernquist_04,Mellema_etal_06}) and a global feature (``reionization step'') in the continuum spectrum of the radio sky that may mark the abrupt overlapping phase of individual intergalactic HII regions \\citep{Shaver_etal_99}.  A key feature for the observation of the line in emission is that the IGM should be heated above the CMB temperature. While pre-heating can happen due to, e.g., dark matter particles decay or annihilation (e.g. \\citealt{MapelliFerraraPierpaoli_06, Valdes_etal_07}), the main source of heating are Ly$\\alpha$ and X-ray photons. Although the impact of these sources has been investigated by several authors (e.g. \\citealt{MadauMeiksinRees_97,ChenMiralda-Escude_04,Nusser_05a,KuhlenMadauMontgomery_06, CiardiSalvaterra_07,PelupessyDiMatteoCiardi_07,RipamontiMapelliZaroubi_08}), none of the above studies has approached the problem considering both sources within a self-consistent model.  Here, we estimate the Ly$\\alpha$ and X-ray photon production using the simulations by \\citet{PelupessyDiMatteoCiardi_07}, which were designed to follow the evolution of both quasar and stellar type sources, including their associated feedback effects. The above estimates are then used to calculate with a semi-analytic approach the evolution of the temperature of an IGM at the mean density.  The paper is structured as follows. In Section~\\ref{sec:simul} we describe the main characteristics of the simulations by \\citet{PelupessyDiMatteoCiardi_07}, while in Section~\\ref{sec:phot_count} we give the prescription used to derive the Ly$\\alpha$ and X-ray photon count. In Section~\\ref{sec:heating} we estimate the IGM heating due to the emitted photons, in Section~\\ref{sec:21cm} the consequences for the observability of the 21~cm line from neutral hydrogen and in Section~\\ref{sec:summary} we give our conclusions. ",
        "conclusions": "\\label{sec:summary}  In this paper we have estimated the effect of X-ray and Ly$\\alpha$ photons on the IGM temperature. This was done calculating the photon production from a population of both stars and quasars based on the self-consistent modeling of these sources and associated feedback processes in the simulations by \\citet{PelupessyDiMatteoCiardi_07}. Such photon production has then been used to determine the evolution of the IGM temperature with a semi-analytic approach. The calculations assume a gas at the mean density whose ionization history is regulated by the effect of the X-ray photons. Thus, they are strictly valid as long as the filling factor of gas ionized by UV photons is small or in regions of the IGM that have not been reached by UV-ionizing radiation.  For reference, the filling factor of ionized regions, under the assumption that it is independent from the effect of the X-rays, assuming a gas clumping factor of 10 and an escape fraction of ionizing photons from stellar type sources of 10\\% is $\\sim 0.004, 0.07, 1$ at $z=15, 10, 7.1$. It should be noted here that the bulk of the UV photons is produced by stars rather than BHs, while the contribution to the ionization fraction from X-rays is only a few percents.  The results presented of course depend quite strongly on the parameters and assumptions adopted in the simulations described in Section~\\ref{sec:simul}. As already mentioned, while e.g. the value of the spin parameter has a negligible effect on the final outcome, the presence of AGN feedback (but not the specific value of $\\epsilon_f$, as long as $\\ne 0$) is crucial. In our reference run we have used a vale $\\epsilon_f=0.5$, which corresponds to coupling only $5\\%$ of the radiated energy. In \\citet{PelupessyDiMatteoCiardi_07} we have shown that no significant difference in the BH accretion rates are seen for $\\epsilon_f=0.5$, but that of course cases with $\\epsilon_f=0$ can imply much larger accretion rates (and associated X-ray photons).  For the sake of illustration, we have performed the same calculations also for the extreme case without feedback, i.e. for $\\epsilon_f=0$. Note that for our calculation this is a completely unrealistic model, which we consider only as a parametric study. The result is shown in Figure~\\ref{fig:temp_igm} as long-dashed line. As expected, the value of the IGM temperature raises much more quickly than in the reference case and becomes larger than $T_{\\rm CMB}$ already at $z \\sim 16$. This would result also in an earlier 21~cm signal in emission.  This case is both unrealistic and also results in a contribution to the unresolved X-ray background which is orders of magnitudes larger than the available observational limits and thus can be discarded.  So it remains that our results depend most strongly on the choice of $M_{h,min}$ and $M_{h,max}$ in eqs.~\\ref{eq:jBH} and~\\ref{eq:sfr}, which in our case we have taken from the masses of the simulated halos. While larger values of $M_{h,max}$ are expected to have a negligible impact because of the paucity of such objects at the redshifts of interest, a smaller $M_{h,min}$ (and its associated high abundance) would result in higher photon production. While extending eq.~\\ref{eq:sfr} to smaller masses is straightforward, it may not be appropriate as the contribution from these low mass halos depends on more complicated forms and strength of the feedback active in the relevant range of redshift (for a review on feedback see \\citealt{CiardiFerrara_05}). Although a consensus on the detailed effects of feedback (in particular radiative feedback) on small structure formation has not been reached yet, most authors agree that the presence of feeback delayes or suppresses the formation of small mass objects, and thus associated star formation, depending on a variety of physical quantities such as redshift, halo mass and molecular hydrogen content (see e.g. \\citealt{MachacekBryanAbel_01,SusaUmemura_06,OSheaNorman_07,Whalen_etal_08a}). A proper treatment of such feedback though is beyond the scope of the present paper.  Similarly the emission properties of mini-QSOs hosted in such small halos is largely unknown and may not be described by the average AGN spectrum here adopted. Thus, we have neglected the contribution of these objects to the total X-ray photon production (see \\citealt{Madau_etal_04, Dijkstra_etal_04, Salvaterra_etal_05} for a discussion about the effect of mini-QSOs in the early universe).  As a reference though, an upper limit for the contribution of Ly$\\alpha$ photon to the IGM temperature is plotted in Figure~\\ref{fig:temp_igm}. Here (long-dashed-dotted cyan line) we show the effects accounting for the contribution to the star formation rate of objects with mass as small as $10^7$~M$_\\odot$. As expected, in this case, the IGM temperature starts to raise at an earlier time due to Ly$\\alpha$ heating and becomes larger than $T_{\\rm CMB}$ at $z \\sim 12$, while X-ray photons dominate heating at $z \\simlt 10$. Note however, this considers that these halos will not produce any X-ray photons, i.e. the accretion onto their black holes must be completely quenched. This is likely a conflicting requirement, as both star formation and black hole accretion will somewhat depend on the same gas supply. To summarize this point, in order to have 21~cm line in emission, X-ray heating from BHs and/or extremely efficient star formation (albeit, as discussed above, somewhat unfeasible) in halos with $M_{h} < 10^8 \\Msun$ is required.  The heating due to stars only in larger halos as shown in Figure~\\ref{fig:temp_igm}, is however unable to bring the gas temperature above $T_{\\rm CMB}$ before the reionization is completed.  To summarize the main results of this work, we find that: \\begin{itemize} \\item unless the Ly$\\alpha$ photon contribution from stars in small mass objects is somehow strongly enhanced in comparison to the X-ray contribution from BHs in the same objects, the IGM heating is always dominated by X-rays;  \\item in this case, the transition redshift between a 21~cm signal in absorption and in emission could be used to constraint BHs accretion properties; \\item for a case consistent with the available limits on the unresolved fraction of the X-ray background in the 0.5-2~keV and 2-8~keV bands, the IGM temperature becomes larger than the CMB temperature (and thus an associated 21~cm signal should be expected in emission) at $z \\simlt 11.5$. \\end{itemize}"
    },
    "0910/0910.4965_arXiv.txt": {
        "abstract": "We present the optical data for 195 HI-selected galaxies that fall within both the Sloan Digital Sky Survey (SDSS) and the Parkes Equatorial Survey (ES).  The photometric quantities have been independently recomputed for our sample using a new photometric pipeline optimized for large galaxies, thus correcting for SDSS's limited reliability for automatic photometry of angularly large or low surface brightness (LSB) galaxies.  We outline the magnitude of the uncertainty in the SDSS catalog-level photometry and derive a quantitative method for correcting the over-sky subtraction in the SDSS photometric pipeline. The main thrust of this paper is to present the ES/SDSS sample and discuss the methods behind the improved photometry, which will be used in future scientific analysis.  We present the overall optical properties of the sample and briefly compare to a volume-limited, optically-selected sample. Compared to the optically-selected SDSS sample (in the similar volume), HI-selected galaxies are bluer and more luminous (fewer dwarf ellipticals and more star formation).  However, compared to typical SDSS galaxy studies, which have their own selection effect, our sample is bluer, fainter and less massive. ",
        "introduction": "Galaxies in the local universe span a range of star formation histories.  At the beginning of the spectrum are gas-rich, low surface brightness (LSB) galaxies that either have just begun the process of star formation or have been processing their gas with extremely low efficiencies.  At the other end are gas-poor galaxies that typically have formed the bulk of their stars in the distant past and are currently devoid of gas (Roberts 1963; McGaugh \\& de Blok 1997).  Creating a sample that bridges these two regimes requires the union of two different methods of identifying galaxies.  Stars dominate the visible light output of most galaxies, and thus galaxies detected by traditional optical imaging have well developed stellar populations. In contrast, the natural way to identify gas-rich, less evolved galaxies is by their 21 cm radio emission. Selecting on HI reveals galaxies entirely based on their gas content, independent of their starlight, and thus easily finds systems with intact gas reservoirs (e.g. Rosenberg \\& Schneider 2000).  Characterizing the stellar and gaseous properties of galaxies selected in the radio and in the optical will therefore yield information about the entire continuum of galaxy star formation histories (Burkholder, Impey \\& Sprayberry 2001).  Aside from its importance for global star formation, a sample of gaseous and stellar information for galaxies in the nearby universe also allows for a more complete census of the local baryons (Rosenberg, Schneider \\& Posson-Brown 2005).  Too often, galaxy studies neglect the fact that HI dominates the baryonic content of many galaxies, particularly those with low masses.  The baryonic makeup of the nearby universe puts important observational constraints on simulations of galaxy formation and evolution as well as revealing reservoirs of mass that were previously undetected because of their optical LSB nature (Disney 1976).  The HI data also provide kinematic constraints on the dark matter content of the galaxies (Blanton, Geha \\& West 2008).  Therefore, with both optical and HI information, we can probe how the baryonic content relates to total mass of the galaxy.  We present the first step toward an inventory of the HI and optical properties of nearby galaxies.  Our study focuses on an HI-selected sample and therefore identifies many systems that have retained much of their primordial HI.  It lacks the systems that have used their entire gas supply and are dominated by stars.  A separate project is underway to complete the nearby baryonic census by filling in the gas poor systems with HI observations of optically selected galaxies.  This study combines data from two high quality, uniform surveys, the Parkes Equatorial Survey (ES; Garcia-Appadoo et al. 2009), and the Sloan Digital Sky Survey (SDSS; York et al. 2000; Gunn et al. 1998; Fukugita et al. 1996; Hogg et al. 2001; Smith et al. 2002; Stoughton et al. 2002; Pier et al. 2003; Ivezi{\\'c} et al. 2004; Gunn et al. 2006; Tucker et al. 2006). The combination of these two surveys creates a rich compendium of stellar and HI parameters for galaxies at various evolutionary states that will be used in future papers to explore how global star formation proceeds in galaxies as a function of their physical parameters.  Previously, most HI surveys were targeted at specific locations, thus little was known about the distribution of HI in the Universe, independent of optical properties. Henning (1992; 1995) used the NRAO 300 ft telescope to conduct an HI blind survey and recovered 39 sources.  Large blind surveys followed using the Arecibo 300m telescope, yielding hundreds of sources and allowing for the first statistically sound studies of the HI mass function (Zwaan et al. 1997; Spitzak \\& Schneider 1998; Rosenberg \\& Schneider 2000; Rosenberg \\& Schneider 2002).  These blind surveys also identified many un-cataloged LSB galaxies and paved the way for more complete studies of the baryonic content of nearby galaxies (Rosenberg, Schneider \\& Posson-Brown 2005).  Recent studies have combined large HI surveys with optical and infrared samples, namely the Arecibo Duel Beam and Slice Surveys with the Two Micron All Sky Survey (2MASS; Jarrett et al. 2000; Rosenberg et al. 2005), the merging of the HI Parkes All Sky Survey (HIPASS) with SuperCOSMOS (Hambly et al. 2001a; Hambly, Irwin \\& MacGillivary 2001b; Hambly et al. 2001c; Doyle et al. 2005), and the combination of HIPASS with DSS/POSSII data (Wong et al. 2009). Rosenberg et al. (2005) were able to probe the baryonic content of a large sample of galaxies, but were limited by the shallow depth of 2MASS, which does not have data for many of the LSB galaxies in the sample. The HIPASS/SuperCOSMOS/DSS samples of Doyle et al. (2005) and Wong et al. (2009) contains optical/IR data for several thousand HI selected galaxies but also suffer from the shallow depth of the SuperCOSMOS/DSS optical data.   \\subsection{The Need for a Uniform Sample} Many studies have investigated the relationships between gas and stars in galaxies (e.g., Roberts 1963; Fisher \\& Tully 1981; Scodeggio \\& Gavazzi 1993; Kennicutt, Tablyn \\& Condon 1994; McGaugh \\& de Blok 1997; Haynes et al. 1999; Burkholder, Impey \\& Sprayberry 2001; Swaters et al. 2002; Iglesias-Paramo et al. 2003; Karachentsev et al. 2004; Helmboldt et al. 2004; Rosenberg, et al. 2005; Serra et al. 2007; Walter et al. 2008; Disney et al. 2008; Zhang et al. 2009). However, many of these have relied on small inhomogeneous samples. These studies have been sufficient to establish the broad trend of increasing gas-richness in low mass systems, but they are limited in their ability to constrain more accurate relationships between gas, stars and galaxy mass, as well as the intrinsic scatter in these physical quantities.  The advent of large astronomical surveys allows for unions of these large surveys to yield multi-wavelength, uniform data sets with small systemic errors and large sample sizes (e.g. Salim et al. 2005; Ag{\\\" u}eros et al. 2005; Obri{\\'c} et al. 2006; Covey et al. 2008).  As mentioned above, several studies have combined large HI surveys with large optical and infrared datasets.  In the first of these, Rosenberg et al. (2005) investigated how the infrared stellar light compares to the HI gas emission.  Although Rosenberg et al. (2005) were able to probe the baryonic content of a large sample of galaxies, they were limited by the shallow depth of 2MASS, which does not have data for many of the LSB galaxies in the sample.  Their study therefore excludes the galaxies at the extreme gas-rich end of the evolutionary spectrum.  Another large-scale blind HI survey is the HI Parkes All Sky Survey (HIPASS; Barnes et al. 2001; Meyer et al. 2004; Zwaan et al. 2004; Wong et al. 2006), which has been combined with IR and optical catalogs (Doyle et al. 2005; Wong et al. 2009) that are very similar to our ES/SDSS catalog (see below).  In fact, a large fraction of the ES data is in the HIPASS catalog (see Garcia-Appadoo et al. 2009 for more information).  While the HIPASS studies have much more sky coverage than our ES/SDSS catalog, the depth of the SDSS allows us to probe optical magnitudes several times fainter and recover optical counterparts for all of the HI sources in our footprint.   The union of SDSS and the ES provides the desired uniformity in both the optical and the radio (HI) data, along with remarkable depth and dynamic range of the optical SDSS data.  Although there is only a modest area of SDSS/ES overlap, enough data exist for the construction of a uniform HI-selected catalog that can be used to probe how the baryonic content of galaxies changes as a function of other physical parameters.  This paper is one of several papers utilizing the combined ES/SDSS data.  In this paper we describe the sample selection, discuss the methods by which we derive the optical photometric parameters, and present the optical data. We briefly describe the sample characteristics and compare our sample to an optically selected sample in a similar volume.  Other papers describe the HI data (Garcia-Appadoo et al. 2009) and explore the gas fractions, colors (West et al. 2009), and dynamics of the ES/SDSS galaxies. ",
        "conclusions": "Using our customized software for analyzing large galaxies in the SDSS, we have created an HI-selected catalog using data from the ES and SDSS surveys.  The data are uniformly sampled and have well characterized uncertainties and limitations. Our software addresses the deblending and sky subtraction issues that make SDSS catalog-level photometry of angularly large galaxies in SDSS unreliable.  We specifically discussed the problems with deblending and sky subtraction and quantified their effect on the ES/SDSS sample.  In addition, we derived a hyper-plane that relates the amount of flux removed by the sky subtraction algorithm to the size, magnitude and axis ratio of the galaxy.  The ES/SDSS galaxies span a large range of surface brightnesses, colors and stellar masses and as expected, the optical properties are different from those of most SDSS galaxy samples (mostly due to selection effects), as well as an optically selected sample compared at the sample volume.  Specifically, the HI selection identifies galaxies with lower surface brightness, smaller absolute magnitudes, bluer colors and smaller stellar masses than those used in typical SDSS studies.  HI-selected galaxies (compared to optically selected galaxies in the same volume) are bluer and brighter, the latter due to the abundance of dwarf elliptical galaxies in optically selected samples.  The ES/SDSS sample also shows a slight bias towards face-on systems, a symptom of the HI-selection.  Other papers incorporate the ES HI data and investigate how the SDSS photometric trends relate to the galaxy gas content (Garcia-Appadoo et al. 2009) and attempt to understand the colors of the galaxies (West et al. 2009).  Future work will examine the dynamics of the ES/SDSS sample and investigate the Tully Fisher relation and its scatter."
    },
    "0910/0910.4362_arXiv.txt": {
        "abstract": "We constrain the amplitude of primordial non-Gaussianity in the CMB data taking into account the presence of foreground residuals in the maps. We generalise the needlet bispectrum estimator marginalizing over the amplitudes of thermal dust, free-free and synchrotron templates. We apply our procedure to WMAP 5 year data, finding $\\fnl= 38\\pm 47$  (1 $\\sigma$), while the analysis without marginalization provides $\\fnl= 35\\pm 42$. Splitting the marginalization over each foreground separately, we found that the estimates of $\\fnl$ are positively cross correlated of $17\\%, 12\\%$ with the dust and synchrotron respectively, while a negative cross correlation of about $-10\\%$ is found for the free-free component. ",
        "introduction": "There has been considerable activity recently in constraining the amount of non-Gaussianity present in CMB data.  This is principally motivated by theoretical interest into deviations of primordial fluctuations from Gaussian statistics--- a natural outcome of several implementations of the inflationary scenario, which can be used as a tool to rule out specific models. More mundanely, non-Gaussianity can also be produced by undetected systematics, which may suggest problems in the dataset. It could also indicate that elements are missing in the standard cosmological model; for example, anomalies which have recently been detected on large scales have been put forward as evidence of cosmic defects or an anisotropic universe.  The primordial non-Gaussianity in inflationary models is usually characterised through the introduction of the $\\fnl$ parameter (see e.~g.~ \\cite{Luo1994,Heavens1998,SpergelGoldberg1999,KomatsuSpergel2001}) which sets the amplitude of the non-linear contribution with respect to the leading order: \\begin{equation} \\label{eq:fnl_eq} \\Phi({\\bf x})=\\Phi_{\\rm L}({\\bf x})+\\fnl(\\Phi_{\\rm L}({\\bf x})^2-\\langle\\Phi_{\\rm L}({\\bf x})^2\\rangle) \\end{equation} where $\\Phi({\\bf x})$ is the primordial gravitational potential and $\\Phi_{\\rm L}({\\bf x})$ is a Gaussian field. In the following we will focus on primordial non-Gaussianity of local type described by Eq.~\\ref{eq:fnl_eq}. Different inflationary scenarios predict different values of $f_{\\rm NL}$; while  in the standard slow roll inflationary scenario \\citep{Guth1981,Sato1981,Linde1982,AlbrechtSteinhardt1982}, $f_{\\rm NL}$ it is predicted to be of the order of unity (see also \\citep{Maldacena:2003,Acquaviva:2003}), in several alternative models \\citep{LythWands2002,LindeMukhanov2006,AlabidiLyth2006,Mizuno2008,Khoury2002PhDT,SteinhardtTurok2002,LehnersSteinhardt2008} it can take much larger values (see e.g. \\cite{Bartolo2004NGreview} for a review). Since the release of the first year WMAP data \\citep{Bennett:2003ca, Komatsu:2003}, there has been a drastic reduction of the upper  limits of $\\fnl$. The most recent constraints coming from WMAP data  using different techniques can be found in \\cite{Komatsu2008wmap5}, \\cite{Smith2009fnlFore}, \\cite{YadavWandelt2007}, \\cite{Curto2009NG}, \\cite{Curto2008waveNG}, \\cite{Pietrobon2008NG}, \\cite{Rudjord2009needBis}, \\cite{Rudjord:2009au}, \\cite{Smidt:2009ir} and \\cite{VielvaSanz:2009}. Results on $f_{\\rm NL}$ with suborbital experiments can be found in \\citep{B2kNG,Curto2008Archeops}. Interestingly, some of these \\citep{YadavWandelt2007,Rudjord:2009au} reported a non-null detection of  $\\fnl$ at more than 2$\\sigma$ level in the WMAP data, leading to a debate on the significance of the signal and whether it may have a cosmological origin or be due to residual foreground or instrumental contamination \\citep{Smith2009fnlFore}.  Non-Gaussianity is also expected from unremoved contamination from astrophysical sources, or foregrounds. When component separation techniques are applied to CMB data (e.g.\\ \\cite{Maino:2002FastICA,Tegmark:2003,Bonaldi:2007,Leach:2008, Bernui:2009}), residual foregrounds remain, and while they are subdominant, they can still be a source of non-Gaussianity, and this could be confused with a primordial signature and affect the constraints on $f_{\\rm NL}$.  In this paper, which is complementary to our previous work \\citep{Pietrobon2008NG}, we aim at generalising the needlet bispectrum estimator to the case in which such foreground residuals are present in the data. Another marginalization technique, using the fast cubic estimator on the bispectrum data, can be found in \\cite{Smith2009fnlFore}.  The paper is organised as follows. In section \\ref{sect:formalism} we give a brief introduction of the needlets and their bispectrum; section \\ref{sect:estimator_gen} addresses the generalisation of the needlet bispectrum estimator in presence of foreground residuals, in section \\ref{sect:dataset} we describe the CMB and the foregrounds dataset used and show our results, in Section \\ref{sect:concl} we comment on our findings. ",
        "conclusions": "\\label{sect:concl} In this paper we have presented a procedure to marginalize the residual foregrounds when estimating $\\fnl$ in the needlet bispectrum framework. However it is important to stress that this algorithm does not strictly rely on needlets properties and it can be easily applied to any linear estimator. With the foreground marginalization, we found, for WMAP 5-year data, that the error bars are enlarged by about $10\\%$ with respect to the estimate obtained without marginalizing. Foreground residuals can have different effects when different estimators are used to characterised  primordial non-Gaussianity. Comparing with other foreground analyses, our results seem to go in the direction of \\cite{YadavWandelt2007}, where they argued that foregrounds negatively biased the $\\fnl$. However, \\cite{Smith2009fnlFore} draw other conclusions showing that the sign of the biasing could depend on the choice of weighting in the method of estimating $\\fnl$. Since needlets coefficients are essentially a rebinning of the filtered harmonic coefficients $a_{lm}$, this could explain the different behaviours. This issue probably needs to be examined separately for each test of non-Gaussianity, since different tests act differently on different spaces (e.g. harmonic, pixel, wavelet), suggesting that the influence of foregrounds on estimating $\\fnl$ is not unique.  All this reinforces the argument that a careful analysis with different tests of non-Gaussianity is crucial to discriminate between primordial non-Gaussianity and spurious effects. Our procedure could be improved further by incorporating any covariance which may arise in the foreground subtraction methods, and by fully considering the effects arising from the potential correlations between different foreground maps. Such analyses will be essential for modelling the non-Gaussianity of future experiments like Planck, where the error bars on $\\fnl$ are expected drastically reduced ($\\Delta \\fnl\\sim 3-5$ \\cite{KomatsuSpergel2001,Babich:2004}), making the uncertainties introduced by foregrounds extremely relevant.  \\subsection*{Acknowledgements} We are grateful to Domenico Marinucci and Francesco Piacentini for fruitful discussions. We thank Michele Liguori, Frode Hansen and Sabino Matarrese for providing us with the primordial non-Gaussian maps. The ASI contract LFI activity of Phase2 is acknowledged."
    },
    "0910/0910.1141.txt": {
        "abstract": "We present new observational determinations of the evolution of the 2--10 keV X-ray luminosity function (XLF) of AGN. We utilise data from a number of surveys including both the 2 Ms \\chandra\\ Deep Fields and the AEGIS-X 200 ks survey, enabling accurate measurements of the evolution of the faint end of the XLF. We combine direct, hard X-ray selection and spectroscopic follow-up or photometric redshift estimates at $z< 1.2$ with a rest-frame UV colour pre-selection approach at higher redshifts to avoid biases associated with catastrophic failure of the photometric redshifts. Only robust optical counterparts to X-ray sources are considered using a likelihood ratio matching technique. A Bayesian methodology is developed that considers redshift probability distributions, incorporates selection functions for our high redshift samples, and allows robust comparison of different evolutionary models. We statistically account for X-ray sources without optical counterparts to correct for incompleteness in our samples. We also account for Poissonian effects on the X-ray flux estimates and sensitivities and thus correct for Eddingtion bias. We find that the XLF retains the same shape at all redshifts, but undergoes strong luminosity evolution out to $z\\sim 1$, and an overall negative density evolution with increasing redshift, which thus dominates the evolution at earlier times. We do not find evidence that a Luminosity-Dependent Density Evolution, and the associated flattening of the faint-end slope, is required to describe the evolution of the XLF. We find significantly higher space densities of low-luminosity, high-redshift AGN than in prior studies, and a smaller shift in the peak of the number density to lower redshifts with decreasing luminosity. The total luminosity density of AGN peaks at $z=1.2\\pm0.1$, but there is a mild decline to higher redshifts. We find $> 50$\\% of black hole growth takes place at $z>1$, with around half in $L_\\mathrm{X}<10^{44}$ erg s$^{-1}$ AGN. ",
        "introduction": "\\label{sec:intro}  Active Galactic Nuclei (AGN) are an important constituent of the Universe, playing a vital role in the formation and evolution of galaxies and having a wider influence on the structure of the Universe. Determining the distribution and evolution of AGN accretion activity throughout the history of the Universe, traced by the luminosity function, is essential to constrain models of super-massive black-hole formation and growth, the triggering and fuelling of AGN and their co-evolution with galaxies. This requires large samples of objects spanning a wide range of redshifts and luminosities to accurately determine the shape of the luminosity function and how this evolves over cosmic time.  X-ray surveys provide a highly efficient method of selecting AGN, including unobscured and moderately obscured sources and low-luminosity sources that may not be identified as AGN at optical wavelengths. X-ray emission is relatively unaffected by absorption, particularly at hard X-ray energies ($\\gtrsim 2$ keV) for moderate column densities ($\\mathrm{N_H}\\lesssim 10^{23}$ cm$^{-2}$), and thus provides a direct probe of the accretion activity. Large efforts have been made to obtain follow-up observations of X-ray sources in various X-ray surveys to determine their redshifts, and thus allow the X-ray luminosity function to be determined \\citep[e.g.][]{Barger03c,Szokoly04,Trouille08}. Such studies of the X-ray luminosity function (XLF) have revealed AGN are a strongly evolving population, with the overall distribution shifting to lower luminosities between $z\\sim1$ and the present day. This evolution can be described by a Pure Luminosity Evolution parameterisation \\citep[PLE,][]{Boyle93,Barger05} in which the XLF retains the same shape (a double power-law with a break at a characteristic luminosity, $L_*$), but shifts to lower luminosities as redshift decreases. However, recent data indicate that a more complex, Luminosity Dependent Density Evolution (LDDE) parameterisation is required to describe the evolution of the XLF at both low ($z\\lesssim1$) and high ($z\\approx 1-4$) redshifts \\citep[e.g.][]{Ueda03,LaFranca05,Silverman08,Ebrero09,Yencho09}. In this scheme the shape of the XLF changes with redshift, with the faint-end slope flattening as redshift increases. This evolution is also characterised by a shift in the peak of the space density of AGN towards lower redshifts for lower luminosities.  However, there are remaining uncertainties as to the exact form of the evolution, especially at high redshifts. In part, this is due to the difficulties of accurately measuring the faint end of the XLF, where incomplete and uncertain redshift information is a serious issue. The \\chandra\\ Deep Fields (CDF-N: \\citealt{Alexander03}; CDF-S: \\citealt{Giacconi02}) have been key in studies of the XLF, providing the deepest available X-ray data and substantial programmes of spectroscopic follow-up of the X-ray sources to determine redshifts. However, even in these fields a high fraction ($\\sim 50$\\%) of the X-ray sources remain spectroscopically unidentified, mainly because the optical counterparts of the faintest X-ray sources are also faint ($R\\gtrsim24$), and thus beyond the limits of current instrumentation. This can severely bias determinations of the XLF, especially at high redshifts \\citep[e.g.][]{Barger05}.  One general approach to address this issue is to determine photometric redshifts for the X-ray sources that are too faint for optical spectroscopic follow-up. This significantly improves the completeness of the samples, although many X-ray sources which lack detectable optical counterparts remain unidentified. However, there are considerable uncertainties in such redshift estimates, particularly for the faintest optical counterparts, and a risk of catastrophic failures in the redshift determinations, which is particularly acute for sources at $z\\sim 1-3$. It is also necessary to correct for remaining incompleteness in such samples, which can prove difficult. Additionally, there is a significant risk of associating the incorrect optical counterpart with an X-ray source, due to the high surface density of optical sources at these faint magnitudes, and thus an X-ray source may be assigned completely incorrect redshift information. For example, \\citet{Barger03c} were able to assign optical counterparts to $\\sim 85$\\% of the X-ray sources in the CDF-N, but noted that $\\sim 25$\\% of the optical identifications at $R=24-26$ are expected to be spurious. Such false identifications of faint X-ray sources could significantly bias determinations of the XLF. To reduce these effects many authors  \\citep[e.g.][]{Ueda03,Barger05,Silverman08} set relatively high X-ray flux limits, thus increasing the spectroscopically identified fraction and reducing the numbers of optically-faint or unidentified sources, but this also reduces the sensitivity to the key low-luminosity AGN. \\amend{The small number of fields with sufficiently high spectroscopic completeness also means most results are susceptible to cosmic variance.}  An alternative approach to the incompleteness issue was adopted by \\citet{Aird08} \\citep[see also][]{Nandra05} to determine the XLF at high redshift, where the effects can be most severe. The Lyman-break technique \\citep[e.g.][]{Steidel93,Steidel03} was used to pre-select a sample of objects at \\z3 based on their broadband colours in three filters, and deep X-ray observations were used to identify AGN. This technique provides an incomplete sample \\amend{(missing heavily reddened sources in particular)}, but the optical selection functions are well-defined and were derived using simulations of the optical data. This allowed corrections for the incompleteness to be applied, and thus the faint end of the XLF at \\z3 could be accurately measured. The number density of low-luminosity AGN was at least as high as previous estimates, and there was no evidence for the flattening of the faint-end slope at $z>1$ that is characteristic of LDDE parameterisations. However, this work did not account for the uncertainties in the redshift estimates for sources without spectroscopic confirmation, and did not present an evolutionary model.  In addition to uncertainties in the optical follow-up and redshift estimates, a number of issues arising from the X-ray observations may also affect the determination of the evolution of the XLF. The faintest sources in \\chandra\\ surveys are detected with very small numbers of counts, and thus Poissonian effects can be significant. Derived X-ray fluxes can thus be highly uncertain, and may be underestimated by up to 50\\% due to the Eddington bias \\citep{Laird09}, leading to uncertainties in the luminosity. \\citet{Georgakakis08} presented a new method of determining the sensitivity of \\chandra\\ observations, which fully accounts for the Poissonian nature of the detection limits, and accounted for the flux uncertainties and Eddington bias in the determination of the X-ray number counts. However, these techniques have not been incorporated in previous studies of the XLF.  In this paper we present a new study of the evolution of the XLF, in which we incorporate new deep X-ray observations and develop improved methodologies to address some of the issues regarding optical counterparts, completeness, photometric redshifts and X-ray fluxes described above. Section \\ref{sec:data} describes our data, including our samples of hard X-ray selected AGN and the data used to perform high-redshift colour pre-selection of AGN. We adopt a likelihood ratio method to ensure we assign secure optical counterparts to the faintest X-ray sources. In section \\ref{sec:bayes_xlf} we develop a Bayesian methodology to determine the evolution of the XLF, which accounts for uncertainties in redshift estimates and X-ray fluxes, incorporates the improved sensitivity calculations of \\citet{Georgakakis08}, corrects for remaining incompleteness of the optical identifications and allows a robust comparison of different evolutionary models. Section \\ref{sec:highz} describes how high-redshift rest-frame UV colour-selected samples are incorporated into this methodology \\citep[building on the work of][]{Aird08}, to provide a more robust determination of the evolution at $z\\gtrsim1.2$.  Our results are presented in section \\ref{sec:xlf_evol}, and discussed further in section \\ref{sec:discuss}.  A flat cosmology with $\\Omega_\\Lambda=0.7$ and $h=0.7$ is adopted throughout. ",
        "conclusions": "\\label{sec:discuss}  Our results for the evolution of the faint end of the 2--10 keV XLF should be the most accurate and robust to date. Our work utilises both of the \\chandra\\ deep fields, including the full 2 Ms exposure available for each field. The additional large area of deep (200 ks) X-ray data provided by the AEGIS survey are also essential for an accurate determination of the faint-end slope. Our sensitive point source detection techniques allow us to probe to the greatest depths in these data, yet our Bayesian techniques allow us to fully account for uncertainty in the flux measurements at these faint limits, and account for the Eddington bias. By incorporating photometric redshift estimates we are able to achieve high redshift completeness ($\\sim 75$\\%), but we have accounted for the uncertainty in such redshift determinations. We have corrected for remaining incompleteness in our samples using the relation between optical and X-ray fluxes, which we have constrained based on our observed samples. At high redshifts we use colour pre-selection to constrain the XLF and reduce biases due to the catastrophic failures of photo-$z$ estimates. We have also used a likelihood ratio matching technique to ensure we only include robust optical counterparts to X-ray sources.  \\subsection{Comparison with previous results} \\label{sec:evol_comp}  Our work offers a number of improvements over previous studies. No previous work has considered the errors in  X-ray fluxes and detection threshold, and thus will not account for Eddington bias. Additionally, no prior studies utilising photometric redshifts \\citep[e.g.][]{Barger05,Silverman08,Ebrero09} have accounted for the known uncertainties in such estimates, or the possible effects of catastrophic failures which we believe led to a biased result at high redshifts in our work (see Figure \\ref{fig:LDDE}, section \\ref{sec:results}). %Photometric redshifts utilised by other authors may be more reliable than our own estimates. The commonly used photometric redshift catalogue for the CDF-N, presented by \\citet{Barger03c}, used the same optical data as our work, and the BPZ code, and thus these estimates will be subject to the same uncertainties and potential for catastrophic failures. \\citet{Zheng04} used a range of photo-$z$ codes, incorporated a wider range of optical filters, deep {\\it HST} imaging data, and used the X-ray data as an indication of absorption properties, to provide improved photometric redshift estimates of X-ray sources in the CDF-S (1 Ms data). Although these estimates were found to be accurate to within $\\sim 8$\\%, the authors discussed the remaining uncertainties, particularly for the fainter sources, and the possibility of catastrophic failure. The lack of spectroscopic follow-up makes it difficult to assess the reliability of all photometric redshifts at high-$z$, and the uncertainties should be accounted for in estimates of the XLF.  Our completeness corrections allow us to account for the lack of redshift information for a fraction of our sources, which is particularly important for determining the faint end of the XLF, where the optical counterparts are often too faint to be detected in even the deepest available imaging. Previous authors have often omitted any form of completeness correction, instead setting high X-ray flux limits \\citep[e.g.][]{Ueda03} or presenting maximal XLFs assigning all unidentified sources to a given redshift bin \\citep[e.g.][]{Barger05}. \\citet{Silverman08} did apply a completeness correction, based on the fraction of sources with a given X-ray flux with redshift information, although samples were restricted to objects with bright optical counterparts ($r'<24$) which may introduce additional biases. \\citet{Ebrero09} restricted their analysis to samples with high-redshift identification completeness, and applied small additional corrections based on the fraction of identified sources as a function of X-ray flux, although the completeness was likely to be over estimated in the CDF-S as they did not consider the high fraction of spurious counterparts. A similar approach was adopted by \\citet{Yencho09}, although their samples were restricted to spectroscopic identifications and thus these completeness corrections were large ($\\sim 60$\\%) at the faintest flux levels, and may be subject to additional biases and inaccuracies given the large number of complex factors involved in spectroscopic success rates. We thus believe our approach of considering only robust optical counterparts to X-ray sources and modelling the distribution of $f_\\mathrm{X}/f_\\mathrm{opt}$ ratios, in addition to our colour pre-selection approach at high-redshifts, constitutes a more robust treatment of completeness than in previous work.   In Figure \\ref{fig:sdens} we directly compare the evolution of the space density of AGN predicted by our LADE model and a number of recent studies of the evolution of the XLF, all of which concluded that LDDE was taking place.  Results are shown in two luminosity bins, which approximately correspond to above and below the characteristic break luminosity, $L_*$, although $L_*$ does evolve and thus such space density plots may provide a distorted view of the evolutionary behaviour of the XLF. Indeed, in the higher luminosity bin ($44.5<\\log \\Lx<46$) our model indicates a very different evolution of the space density at $z\\lesssim 1$ to previous (LDDE) results, but this is partly due to small offsets in luminosity space which manifest as very large differences in density because of the steep slope of the bright end of the XLF. This can also lead to large differences in binned estimates depending on the methodology or assumed evolution of the XLF. However, in this regime, the XLF remains poorly constrained, and despite the very different evolutionary behaviour our binned results are consistent with the \\citet{Silverman08} and \\citet{Ebrero09} models at the $\\sim 1-2\\sigma$ level, although the \\citet{Yencho09} model is significantly lower. The high redshift form of our evolution in this luminosity bin predicts similar space densities to the previous studies, although the form of the evolution and location of the peak in the number density is somewhat different.  Our results are in good agreement at low redshifts ($z\\lesssim1$) in the lower luminosity bin ($43<\\log \\Lx<44.5$), although the \\citet{Yencho09} results remain below our estimates. However, at higher redshifts our model predicts a much higher space density than found by \\citet{Yencho09} or \\citet{Silverman08}. We attribute the decline found by these authors to incompleteness, catastrophic failures of photometric redshift estimates, and false associations of X-ray sources with lower redshift optical counterparts. Conversely, the \\citet{Ebrero09} LDDE model is in very good agreement with our result, most likely because of the higher completeness of their samples. However, our results have shown that an LDDE model is not necessary to describe this evolution.  It is worth noting that our LADE form of the evolution is able to describe a ``cosmic-downsizing\" behaviour in which the peak in the space density moves to lower redshifts for lower luminosity AGN, and indeed we do find small redshift offsets between the peak space densities for different luminosity ranges (e.g. Figure \\ref{fig:sdens}). However, our model does not find the strongly luminosity-dependent shift in peak space density indicated by some prior studies \\citep[e.g.][]{Hasinger05,Silverman08}. This is mainly due to our steeper faint-end slope (and thus higher number density of low-luminosity AGN at $z\\gtrsim 1$), but is also due to the differences in the form of the bright-end evolution discussed above.   % Space density with comparisons \\begin{figure} \\includegraphics[width=\\columnwidth]{f10_col.eps} \\caption{ The evolution of the space density of AGN based on our LADE model for two luminosity ranges: $43<\\log \\Lx<44.5$ (\\textit{circles}) and $44.5<\\log \\Lx < 46$ (\\textit{squares}). Binned estimates are shown from the hard X-ray samples at $z<1.2$ (\\textit{solid symbols}) and the colour pre-selected samples at \\z2--3 (\\textit{open symbols}). The space density based on LDDE models determined by \\citet[\\textit{dot-dashed line}]{Ebrero09}, \\citet[\\textit{dashed line}]{Yencho09} and \\citet[\\textit{dotted line}]{Silverman08} are also shown. } \\label{fig:sdens} \\end{figure} %%%%%%%%%%%%%%%%%  \\subsection{Remaining uncertainties} \\label{sec:limitations}  While we believe our work offers a number of improvements over previous studies, there are remaining uncertainties associated with our approach.  \\subsubsection{Photometric redshifts} As previously discussed, the introduction of photometric redshifts, essential to improve the completeness of samples at faint X-ray fluxes, does introduce significant uncertainty, which we have tried to account for by considering probability distributions for the redshift estimates. However, there are a number of limitations to photometric redshift techniques that are thus also limitations in our studies. This is particularly apparent in the catastrophic failure of our photo-$z$ estimates at $z\\gtrsim1.2$, which we believe led to a biased result for the evolution of the XLF, favouring an LDDE scheme with a flattening faint-end slope (section \\ref{sec:results}, table \\ref{tab:xlf_evol}). Additionally, the majority of our photo-$z$ estimates are based on template fitting approaches that utilise a variety of galaxy templates (and priors based on the galaxy population), and thus the derived redshift probability distributions may not correspond to the true redshift for an AGN, which could bias our results at low as well as high redshifts. Incorporating AGN templates and using priors based on the AGN population may offer significant improvements and reduce systematic biases and catastrophic failures \\citep[e.g.][]{Zheng04,Salvato08}. Expanding the number and wavelength range of the filters, such as including near- and far-IR data, also has potential to improve photometric redshift estimates and provide better fits to AGN templates \\citep[e.g][]{Salvato08}.  In addition to the template fitting approach, we also utilised ANN$z$ redshift estimates in the AEGIS field, in which an artificial neural network is trained to provide photometric redshifts using a sample of objects with spectroscopy. While this approach produced reliable results, it is limited to the subset of objects with a certain range of colours and magnitudes for which follow-up spectroscopy is available. Indeed, spectroscopic validation of photo-$z$ estimates for large numbers of optically-faint X-ray sources is impossible with current instrumentation, which is a major limitation of photometric redshift techniques.  \\subsubsection{High-$z$ colour pre-selection}  To reduce the bias at high redshifts we incorporated rest-frame UV colour pre-selected samples. This should avoid the systematic biases associated with catastrophic failure of the photometric redshift estimates, but does risk introducing additional uncertainty and biases. As noted in section \\ref{sec:results}, the binned estimates based on the colour pre-selected samples at $2.0<z<2.5$ (dominated by BX selected objects) fell systematically below the model, ostensibly indicating a rise in the space density of AGN between \\z2 and \\z3 (also seen in Figure \\ref{fig:sdens}). The BX criteria select objects in a very thin slice of colour space, identifying objects with a specific spectral shape rather than strong spectral features. This makes the BX selection both more susceptible to contamination by lower redshift sources and more prone to incompleteness than the more widely-used LBG selection that identifies a strong break in the spectrum shortward of the Lyman-limit. The shape and efficiency of the derived selection functions will be highly dependent on the modelling of the AGN spectra and the scatter in their properties, and it is quite possible that our relatively simple modelling may have led to overestimates of the completeness. Indeed, both the LBG and BX selections require AGN to be bright in the rest-frame UV, and will therefore miss low-luminosity, moderately or heavily obscured AGN residing in red, evolved host galaxies. \\citet{Aird08} found that LBG selection recovered a number density of low-luminosity AGN that is at least as high as found via direct X-ray selection and follow-up, tentatively indicating that the AGN population at \\z3 is dominated by objects with blue colours \\citep[cf. \\z1,][]{Nandra07b}. However, this may not be true at slightly lower redshifts, resulting in a lower number density at $2.0<z<2.5$ via the BX selection as red galaxies will be missed.  As discussed in section \\ref{sec:lowzinterlopers}, our combination of LBG and X-ray selection to identify \\z3 AGN does not suffer from low-redshift contamination, but this is not the case for the BX selection. Without complete spectroscopic follow-up of the X-ray detected BX candidates we cannot be confident that we have excluded all interlopers (although we believe we are able to identify low-redshift star-forming galaxies, see Figure \\ref{fig:fxfopt_highz}), and it is unclear how this might affect our derived evolution of the XLF. Extensions of the BX selection technique, which incorporate data at redder wavelengths, may  allow low-redshift interlopers to be excluded, and help identify AGN in redder host galaxies. Indeed, a hybrid of the photometric redshift and colour pre-selection approaches, utilising data from all available wavebands to perform SED fitting and estimate a redshift, but also performing a number of colour cuts to avoid catastrophic failures, and carefully modelling the distribution of AGN spectral properties to determine the incompleteness, may be the optimal approach to determine the high-redshift evolution.  It is also worth noting that our approach of combining direct hard X-ray selection at $z<1.2$ and colour pre-selection at high redshifts provides no data for $z\\approx 1.2-2.0$. Significant evolution may be taking place over this redshift range, particularly at high luminosities \\citep[e.g.][]{Silverman05}. Indeed, at high luminosities incompleteness issues should be less severe, although we note that at \\z2 we expect $\\sim$10\\% of AGN with $\\Lx=10^{44.5}$ \\ergs\\ will have optical magnitudes $R>24$, based on our $\\fx/f_\\mathrm{opt}$ distributions. In this redshift range, photometric redshift estimates are prone to catastrophic failure, and spectroscopic identification can also be difficult even for bright targets. Thus care must be taken to avoid biased results. As found in section \\ref{sec:results}, such systematic biases can erroneously indicate the need for a more complex evolutionary model.   \\subsubsection{Completeness corrections}  To account for objects with optical magnitudes below our limit of $R=25.5$, we introduced a simple completeness correction assuming the $\\log \\fx/f_\\mathrm{opt}$ distribution can be described by the same Gaussian function at all redshifts (section \\ref{sec:completecorr}). This may be an over simplification, as the redshift distribution of optically-faint sources will differ from the brighter population. Indeed, \\citet{Mainieri05} found that optically-faint X-ray sources tended to be found at higher redshifts. Additionally, \\citet{Koekemoer04} discuss how Extreme X-ray Objects (EXOs) which are X-ray bright but lack optical counterparts, even in very deep \\textit{Hubble} ACS imaging, may represent a population of very high-redshift ($z\\sim 6-7$) AGN where the redshifted Lyman-break suppresses all emission at optical wavelengths. More sophisticated modelling of the rest-frame UV--optical properties of the AGN population and the effects on observed magnitudes as a function of redshift would enable improvements to our completeness corrections. Furthermore, the shape and evolution of the optical luminosity function of the AGN host galaxies will affect the completeness of our samples , could alter the shape and efficiency of the optical selection functions, and may systematically impact the photo-$z$ estimates. In addition, we do not currently utilise any information from the optically-unidentified population (their X-ray fluxes, the fact that they lack counterparts to the limits of our imaging), which could place further constraints on the evolution of the XLF.   \\subsubsection{Cosmic variance} Our investigation of the XLF uses an unprecedented large area of deep 200 ks \\chandra\\ data in the AEGIS field, which supplements the data from the \\chandra\\ Deep Fields used in previous studies \\citep[e.g.][]{Barger05,Silverman08}. As such, our determination of the faint-end evolution should be less susceptible to the effects of cosmic variance than in prior work. However, cosmic variance could introduce additional uncertainty in our XLF determinations (and thus weaken evidence for more complex evolutionary schemes). Indeed, assuming our X-ray detected AGN cluster according to a power-law with scale length $r_0=7.1$ Mpc and slope $\\gamma=1.8$ \\citep[typical of X-ray AGN at $z\\sim1$][]{Coil09}, we predict a fractional RMS variation of $\\sim 5-10$\\% in the number density of AGN. \\amend{This is comparable to our Poissonian uncertainties due to the number of objects in bin, although the errors will be basically covariant across the luminosity bins, and thus will not affect the shape of the XLF.}  %We do not know if such trends will extend to the unidentified population, but we can certainly not rule out the possibility they have a very different redshift distribution to the observed sources, and this could change our conclusions regarding the evolution of the XLF. We have also made no attempt to include the information we \\emph{do} have about these sources in our determination of the XLF: we have measurements of their X-ray fluxes, and we know that their optical counterparts are likely to be fainter than $R\\approx 25.5$. Incorporating this information directly is a potential extension of our work.  \\subsubsection{Intrinsic absorption}  We have calculated rest-frame 2--10 keV luminosities for our sources from their hard band fluxes (or soft band for the high-$z$ sample), assuming a power-law with $\\Gamma=1.9$. This approach should correct for absorption effects due to intrinsic column densities of $\\mathrm{N_H}\\lesssim10^{23}$ cm$^{-2}$ at $z\\sim1$. However, the effectiveness of this correction varies with redshift, and this may bias our measurements of the luminosity, and alter the effective sensitivity of our observations. The extent of these effects will depend on the distribution of AGN absorption properties, which may be evolving with redshift or have a luminosity dependence \\citep{Ueda03,LaFranca05}. Some previous studies \\citep{Ueda03,LaFranca05,Ebrero09} instead attempted to correct for the effects of absorption on a source-by-source basis, and directly account for the $\\mathrm{N_H}$ distribution in their calculation of the XLF. Determining the extent of absorption is difficult for our faint sources, which have few counts and thus little X-ray spectral information. The effects of absorption on our X-ray sensitivity and completeness calculations will be complicated, and should be corrected for a number of potential selection biases . Such effects are in fact likely to introduce considerable additional uncertainty as to the evolution of the faint end of the XLF, and could thus strengthen our assertions of the lack of a requirement for a more complicated LDDE model to describe the evolution of the XLF.  We defer these considerations to future work.  %%%%%% % \\subsection{The evolution of AGN accretion activity} \\label{sec:conc_evol}  Our determination of the evolution of the XLF presented in this paper sheds new light on how the distribution of AGN accretion activity evolves over the history of the Universe. We conclude that given the currently available data, the evolution of the XLF can be described by a luminosity and density evolutionary scheme, in which the XLF retains the same shape, but shifts in both luminosity and density with redshift. A more complicated luminosity-dependent scheme does not provide a significantly better description of the data; even so our best-fit LDDE for the hard X-ray sample only, which we expect is biased at high redshifts, indicates a steeper faint-end slope and less evolution of the shape of the XLF than most previous work.  \\citet{Hopkins05} presented an interpretation of the AGN luminosity function in which the bright end traces the mass distribution of black-holes accreting at the peak rate over their lifetimes, while the faint end corresponds to AGN during transitionary periods as they approach or decline from their peak rate of activity. Furthermore, \\citet{Hopkins06} proposed that a luminosity-dependent lifetime for AGN could explain any observed flattening of the faint-end slope, as less time is spent in the transitionary period by the more luminous AGN found at high redshifts. Our results, however, show that the shape of the XLF does not change, thus a strongly luminosity-dependent lifetime may not be required. The lack of evolution in the shape of the XLF indicates the processes involved in the triggering and feeding of AGN remain the same at all redshifts, but increase in overall density from the earliest times until \\z1, and moving to lower luminosity systems between $z\\approx 1$ and the present day. Such behaviour could be consistent with a merger-driven model of the build-up of galaxy and black-hole mass, and the triggering of AGN \\citep[e.g.][]{Kauffmann00,Hopkins08}, in which the hierarchical build-up of galaxies increases the number density of AGN in the early history of the Universe, but the exhaustion of gas supplies in the most massive galaxies at later times \\citep[e.g.][]{Cowie96} results in the ``down-sizing\" of AGN to lower luminosities. Fully reconciling such a picture with the full range of observational data for AGN, in addition to the observed evolution of galaxies, using current theoretical models clearly requires significant further investigation. Indeed a number of other processes may play a role in triggering and fuelling low luminosity AGN, such as minor mergers \\citep{Hernquist95}, accretion of gas from a hot halo \\citep{Fabian95} or bar-driven accretion \\citep{Sellwood99}. It is essential to assess the importance of such processes and understand the role they play in the evolution of AGN accretion activity. %%%%%%%%%% % ldens \\begin{figure} \\includegraphics[width=\\columnwidth]{f11_col.eps} \\caption[Space density and luminosity density of AGN as a function of redshift]{ 2--10 keV luminosity density of AGN as a function of redshift for our LADE model integrated over the luminosity ranges indicated. The grey shaded region indicates the 1$\\sigma$ uncertainty on the total luminosity density of AGN. The luminosity density is dominated by moderate luminosity ($43<\\log \\Lx<45$) AGN, and exhibits a mild decline above $z\\approx 1.2$.} \\label{fig:ldens} \\end{figure} %%%%%%%%%%  In Figure \\ref{fig:ldens} we plot the 2--10 keV luminosity density, $\\int \\Lx \\phi(\\Lx,z) \\dd \\log \\Lx$, as a function of redshift,  which provides a tracer of the total AGN accretion activity. The evolution of the total luminosity density for all AGN with $\\Lx>10^{42}$ \\ergs\\ is shown as well as the contribution of AGN in set luminosity ranges. The total luminosity density peaks at $z=1.2\\pm0.1$, rapidly declining to lower redshifts, but the decline to higher redshifts is much milder indicating significant AGN activity is taking place out to \\z3--4. AGN with luminosities in the range $\\Lx\\approx 10^{43-45}$ \\ergs (i.e.  $\\sim L_*$) are responsible for the majority of the luminosity density, although the contribution of $\\Lx=10^{44-45}$ \\ergs\\ AGN falls off rapidly at $z\\lesssim 0.8$ as $L_*$ evolves, and $\\Lx=10^{43-44}$ \\ergs\\ AGN dominate.  We can track the build-up of black-hole mass more directly by relating AGN luminosity to mass accretion, as first proposed by \\citet{Soltan82}, \\begin{equation} L_\\mathrm{bol}=\\epsilon \\dot{M_\\mathrm{acc}}c^2= \\frac{\\epsilon \\dot{M_\\mathrm{bh}}c^2}{1-\\epsilon} \\label{eq:macc} \\end{equation} where $\\dot{M_\\mathrm{acc}}$ is the mass accretion rate, $\\dot{M_\\mathrm{bh}}$ is the rate of change of black-hole mass density, $L_\\mathrm{bol}$ is the bolometric luminosity, $\\epsilon$ is the radiative efficiency of the accretion process and $c$ is the speed of light. We have adopted a simple approach \\citep[e.g.][]{Barger05}, converting our 2--10 keV X-ray luminosities to bolometric values using a constant conversion factor \\citep[40,][]{Elvis94}, and assuming a single value of the radiative efficiency, $\\epsilon=0.1$ \\citep{Marconi04}. We note that a number of authors have discussed the need for luminosity-dependent bolometric corrections \\citep[e.g.][]{Marconi04,Hopkins07b}, and \\citet{Vasudevan07} reported significant object to object variation in bolometric corrections, which may depend on the Eddington ratio. The radiative efficiency may also vary significantly between AGN, depending on the spin of the black-hole \\citep[e.g.][]{Thorne74} in addition to the specifics of the accretion processes \\citep[e.g.][]{Merloni08}. However our simple assumptions allow an initial investigation of the consequences of our derived XLF evolution on the build-up of black-hole mass. Using equation \\ref{eq:macc} we can convert our luminosity densities to a mass accretion rate, and thus calculate the total black-hole mass density as built up by accretion activity over the history of the Universe. Our results are shown in Figure \\ref{fig:bh_acc}.  %%%%%%%%%% % BH mass density \\begin{figure} \\begin{center} \\includegraphics[width=\\columnwidth]{f12_col.eps} \\end{center} \\caption[Accreted black-hole mass density against lookback time]{ Total accreted black-hole mass density against lookback time based on our LADE ({\\it solid}, grey shaded region indicates 1$\\sigma$ uncertainty in the derived model). The lower curves correspond to the same luminosity ranges indicated in Figure \\ref{fig:ldens}. We find $\\sim 50$\\% of the local black-hole mass density is built up in AGN actively accreting at $z\\gtrsim1$. } \\label{fig:bh_acc} \\end{figure} %%%%%%%%%%%  Based on our LADE model we predict a local black-hole mass density of $2.2\\pm0.2 \\times 10^5\\; M_{\\sun}$ Mpc$^{-3}$, where the error reflects the uncertainties in our model fit, but not the potentially larger uncertainties in bolometric correction or accretion efficiency. This value is in good agreement with the estimate of \\citet{Yu02} ($2.5\\pm0.4 \\times 10^5\\; M_{\\sun}$ Mpc$^{-3}$) based on velocity dispersions of early-type galaxies in SDSS and the $M_\\mathrm{bh}-\\sigma$ relation, although is lower than the estimate of \\citet{Marconi04} ($4.6^{+1.9}_{-1.4} \\times 10^5\\; M_{\\sun}$ Mpc$^{-3}$), possibly indicating that our XLF does not provide a complete census of the history of accretion activity. Figure \\ref{fig:bh_acc} shows that a significant fraction ($\\sim 50$\\%) of this total mass density is accreted at $z\\gtrsim 1$. While the majority of the mass build up takes place in moderate luminosity AGN ($\\Lx=10^{43-45}$ \\ergs), a significant fraction is accumulated at lower luminosities ($\\Lx=10^{42-43}$ \\ergs). The LDDE model from section \\ref{sec:xlf_evol} predicts a lower local black-hole mass density, mainly due to the smaller numbers of AGN at these low luminosities and high redshifts. The redshift range \\z1--3 clearly corresponds to a period of significant AGN activity, and thus it is essential to accurately measure the XLF down to $\\Lx\\approx 10^{42-43}$ \\ergs\\ in this epoch to determine the history of black-hole mass accretion.  We can also compare our derived black-hole mass accretion rates to star-formation rates, which we show in Figure \\ref{fig:sfr}. Previous studies \\citep[e.g.][]{Boyle98b,Silverman08} have shown close similarities between the rapid increase in star-formation rate and black-hole accretion at $z\\lesssim 1$, which we confirm. Our results indicate that this correlation continues out to high redshifts \\citep[cf.][]{Silverman08}, at least when comparing to the recent star-formation rates of \\citet{Bouwens07}, although we note that our model is extrapolated far beyond the redshift range probed by our data. Comparisons of the galaxy and AGN luminosity functions may reveal differences in the details of the evolving distributions of activity, which could reveal further facets of the co-evolutionary processes and the feedback regulating AGN activity and star-formation. %In particular, we note that the galaxy luminosity function is found to be well described by a Schecter function \\citep[e.g.][]{Efstathiou88,Lin96,Blanton03}, with a steep exponential drop off at high luminosities, rather than the power-law dependence found for the AGN luminosity function. This could imply that accretion and the feeding of black-holes in the highest mass systems is significantly more efficient than star-formation processes, which may be halted by strong feedback from the AGN.  %%%%%%%%%% % SFR comparison \\begin{figure} \\begin{center} \\includegraphics[width=\\columnwidth]{f13_col.eps} \\end{center} \\caption[Comparison of the star-formation and black-hole accretion rates as a function of redshift]{ Comparison of star-formation rates ({circles}) from the compilation of \\citet{Hopkins04} and black-hole accretion rate based on our LADE model ({\\it solid line}; grey shading indicates 1$\\sigma$ uncertainty in the model). Black-hole mass accretion rate assumes a constant bolometric correction and accretion efficiency, $\\epsilon=0.1$, and is scaled up by a factor 5000 \\citep{Silverman08}. We also show recent measurements of the star-formation rate at high redshifts from \\citet[{\\it squares}]{Bouwens07}, which indicate the correlation between star-formation and black-hole accretion rate continues at high redshifts. } \\label{fig:sfr} \\end{figure} %%%%%%%%%  %%%%%%%%% % Summary"
    },
    "0910/0910.2402_arXiv.txt": {
        "abstract": "By analyzing the \\x\\ spectrum of \\mcg\\ obtained with the \\hetgs\\ spectrometer on board the \\chandra\\ observatory, we identify three kinematically distinct absorption systems; two outflow components intrinsic to \\mcg, and one local at $z = 0$. The slow outflow at $-100~\\pm\\ 50$~\\kms\\ has a large range of ionization manifested by absorption from 24 different charge states of Fe, which enables a detailed reconstruction of the absorption measure distribution ($AMD$). This $AMD$ spans five orders of magnitude in ionization parameter: $-1.5 < \\log\\xi < 3.5$ (\\cgs), with a total column density of $N_H$ = (5.3~$\\pm$ 0.7) $\\times10^{21}$~\\cmsq. The fast outflow at $-1900~\\pm 150$~\\kms\\ has a well defined ionization parameter with $\\log \\xi = 3.82\\pm 0.03$ (\\cgs) and column density $N_H~= 8.1~\\pm\\ 0.7 \\times10^{22}$~\\cmsq. Assuming this component is a thin, uniform, spherical shell, it can be estimated to lie within 11 light days of the \\agn\\ center. The third component, most clearly detected in the lower oxygen charge states O$^{+1}$ - O$^{+6}$, has been confused in the past with the fast outflow, but is identified here with local gas ($z = 0$) and a total column density $N_H$ of a few $10^{20}$~\\cmsq. Finally, we exploit the excellent spectral resolution of the \\hetgs\\ and use the present spectrum to determine the rest-frame wavelengths  of oxygen inner-shell lines that were previously uncertain. ",
        "introduction": "\\label{sec:intro}  Approximately half of type 1 \\agn s show complex absorption in the soft \\x\\ band \\citep{crenshaw03}. The blue shifted absorption lines come from a highly or partly ionized outflow first noted by \\citet[]{halpern84}. These outflows may play a central role in cosmological feedback, in the metal enrichment of the intergalactic medium (IGM) and in understanding black hole evolution. However, the outflow physical properties such as mass, energy and momentum are still largely unknown.A necessary step to advance on these issues is to obtain reliable measurements of the ionization distribution, and column densities of the outflowing material.  \\mcg\\ is a bright Seyfert~1 galaxy ($L_X\\sim10^{43}$ erg~s$^{-1}$) at a redshift z~=~0.007749 \\citep{fisher95}. It shows strong optical reddening \\citep{reynolds97}, which suggests that the absorber may also contain some cold (neutral) gas. \\mcg~ is a highly variable \\x\\ source with changes up to a factor of 2 on time scales of $\\sim$ 1 ks \\citep{fabian94,reynolds95,otani96} \\mcg\\ is well known for its relativistically broadened emission feature at 5~-- 7 keV \\citep[]{tanaka95,ba03,young05}  that has been interpreted as Fe K$\\alpha$ emission from the accretion disk deep inside the black hole gravitational potential well.  There are numerous studies in the literature of the ionized \\x\\ absorber of \\mcg\\ the most important of which are listed in Table~\\ref{table1}. \\citet{otani96} analyzed the \\asca\\ data and \\citet{reynolds97} fitted the \\rosat\\ and \\asca\\ spectra using strong, variable, oxygen absorption edges to explain the curved continuum shape. In 2000, \\citet{brand01} obtained the first grating spectra of \\mcg\\ with the Reflection Grating Spectrometer (\\rgs) on board \\xmm. The improved resolution showed the spectral turn overs to be inconsistent with the oxygen edge positions. Indeed, \\citet{brand01} modeled the soft \\x\\ spectrum with relativistic emission lines from deep in the accretion disk of C$^{+5}$, N$^{+6}$ and O$^{+7}$, instead of absorption edges. \\citet{lee01} then used the 2000 \\chandra~ High Energy Transmission Grating (\\hetgs) spectrum of \\mcg\\ to claim that (O$^{+6}$ but mostly) neutral Fe L-shell edges can explain the continuum turn over at $\\approx 17.5$~\\AA\\ \\citep[see also][]{lee09}, and not relativistic emission lines. As in previous works, \\citet{lee01} needed two ionization zones in order to fit the soft \\x\\ absorption lines. \\citet{sako03} reinstated the \\citet{brand01} interpretation of the 2000 \\rgs\\ data, but with a better absorption model. \\citet{sako03} identified two velocity components in the absorber with outflow velocities of --150 \\kms\\ and --1900 \\kms . \\citet{turner04} fitted a second longer 320~ks \\rgs\\ observation from 2001, and confirmed the two kinematic components, with (roughly consistent) velocities of +80~$\\pm$ 260~\\kms\\ and --1970~$\\pm$ 160~\\kms. \\citet{young05} studied a long 520 ks 2004 \\chandra\\ \\hetgs\\ observation, and found some of the most highly ionized species of Fe$^{+24}$, Fe$^{+25}$, S$^{+15}$ and Si$^{+13}$ to be outflowing at --2000 \\kms\\ . The analysis of \\citet{mckernan07} of the previous 2000 \\hetgs\\ spectrum again confirmed the fast component, but with a slightly different velocity of $-1550^{+80}_{-130}$ \\kms. They also reported two slow components with low velocities of $\\leq$~30~\\kms\\ and $\\leq$~15 \\kms. Recently, \\citet{miller08} studied all of the archival spectra of \\mcg\\ including a 2006 \\suzaku\\ observation and confirmed the fast outflow and the slow outflow with its two ionization components. \\citet{miller08} also invoked a partial covering absorber that helps explain the continuum shape without the need of relativistically broadened emission lines. Most recently, \\citet{chelouche08} modeled the 2004 \\hetgs\\ spectrum with similar components as the previous authors. We summarize the long list of previous works on this intriguing target and in particular its grating observations by stating that while the gratings have led to unambiguous measurements of the kinematics and ionization of the absorbing outflow, the physical interpretation of the \\x\\ continuum is still being debated.  In this paper, we wish to further investigate the physical conditions of the \\mcg\\ \\x\\ absorber, using the archival 520 ks \\hetgs\\ observation from 2004, but with special focus on the ionization distribution of the plasma, and on lines that do not seem to fit the simple two-velocity picture \\citep[see list of individual-line velocities in][]{turner04}. While the majority of works on the \\x\\ spectra of \\agn\\ outflows employ gradually increasing number of ionization components, until the fit is satisfactory \\citep[e.g.,][]{kaspi01, sako03}, it is instructive to reconstruct the actual distribution of the column density in the plasma as a continuous function of ionization parameter $\\xi$ \\citep{katrien05}, which we termed the absorption measure distribution \\citep[$AMD$,][]{tomer07}. Although for high-quality spectra such as the present one, two or three ionization components might produce a satisfactory fit, the $AMD$ reconstruction is the only method that reveals the actual distribution including its physical discontinuities (e.g., due to thermal instability), and ultimately provides a more precise measurement of the total column density. ",
        "conclusions": "\\label{sec:concl}  We have analyzed the kinematic and thermal structure of the ionized outflow in \\mcg. We find three distinct absorption systems, two of which are intrinsic to the \\agn. The slow component is outflowing at --100~\\kms\\ and spans a considerable range of ionization from neutral Fe to Fe$^{+23}$ ($-1.5 < \\log \\xi< 3.5$~\\cgs). A second, fast outflow component at --1900~\\kms\\ is very highly ionized ($\\log \\xi = 3.8$~\\cgs). Finally, a third component of local absorption at $z = 0$ is detected for the first time in \\mcg\\ by its oxygen absorption.  Using our $AMD$ reconstruction method for the slow component, we measured the distribution of column density as a function of $\\xi$. We find a double-peaked distribution with a significant minimum at $0.5 < \\log \\xi < 1.5$ (\\cgs) , which corresponds to temperatures of $4.5 < \\log T < 5$ (K). This minimum was observed in several other \\agn\\ outflows and it can be ascribed to thermal instability that appear to exist ubiquitously in photo-ionized Seyfert winds. The fast outflow with its narrow ionization distribution can be described as a thin shell and estimated to be approximately ten light days away from the central ionizing source. The local absorption system we believe could arise from either the ionized Galactic ISM, or from the Local Group. Finally, we use the \\hetgs\\ spectrum to slightly improve on the computed wavelengths of the most important inner-shell O lines."
    },
    "0910/0910.4859_arXiv.txt": {
        "abstract": "Magnetars are young neutron stars with extreme magnetic fields (B$\\ga$10$^{14}$-10$^{15}$\\,G). How these fields relate to the properties of their progenitor stars is not yet clearly established. However, from the few objects associated with young clusters it has been possible to estimate the initial masses of the progenitors, with results indicating that a very massive progenitor star ($M_{\\rm prog} >$40\\msun) is required to produce a magnetar. Here we present adaptive-optics assisted Keck/NIRC2 imaging and Keck/NIRSPEC spectroscopy of the cluster associated with the magnetar SGR~1900+14, and report that the initial progenitor star mass of the magnetar was a factor of two lower than this limit, $M_{\\rm prog}$=17$\\pm$2\\msun. Our result presents a strong challenge to the concept that magnetars can only result from very massive progenitors. Instead, we favour a mechanism which is dependent on more than just initial stellar mass for the production of these extreme magnetic fields, such as the ``fossil-field'' model or a process involving close binary evolution. ",
        "introduction": "\\label{sec:intro} Magnetars are currently recognized as a distinct group of neutron stars comprising of several classes of object, such as Soft Gamma Repeaters (SGRs), Anomalous X-ray Pulsars (AXPs), and some Compact Central Objects. These objects are characterized by relatively long spin periods and large spin-down torques, which imply magnetic fields on the order of B$\\ga10^{14}-10^{15}$\\,G, making them the most highly magnetized objects known in the Universe \\citep{D-T92,T-D95,Kouveliotou98,Mereghetti08}. They are also known to enter active periods during which they emit very intense ($10^{37}\\la L \\la10^{41}$ erg/s), short ($\\sim$0.1\\,s) hard X-/gamma ray bursts, as well as extremely energetic Giant Flares of $\\ga10^{44-46}$ ergs lasting several minutes.  It is still unclear how magnetars are formed. The current theoretical framework for magnetar production requires that the core of a massive star has a very fast rotation speed in the first few seconds after it goes supernova (SN). If the rotation period is shorter than the convective timescale within the neutron star -- about 1ms -- a highly-efficient dynamo operates which boosts the magnetic field to $\\sim$1000 times that of a `regular' neutron star, and very rapidly slows the rotation period down to of order 1 second in a matter of years \\citep[the so-called `dynamo' mechanism,][]{D-T92,TCQ04}. However, recent stellar evolution calculations show that the cores of massive stars are substantially spun down as they enter the Red Supergiant (RSG) phase through magnetic braking between the stellar core and convective envelope \\citep{Heger05}. Thus, the problem exists of how the core of a massive star can retain sufficient angular momentum through to the SN stage such that the post-SN core is able to jump-start the dynamo mechanism. It has been suggested that those stars with $M_{\\rm init} \\ga $40\\msun\\ are able to lose a substantial fraction of their hydrogen-rich envelope while still on the main-sequence, allowing them to skip the RSG phase, and therefore avoid the severe spin-down of the core as the outer envelope expands and becomes convective \\citep{Gaensler05apj}.  Where magnetars have been associated with star-clusters it has been possible to estimate the initial mass of the progenitor empirically. Evidence suggests that the magnetar phase is short and that the SN that produced it must have occured recently \\citep[$\\la10^4$yrs ago,][]{Kouveliotou94}, such that the age of the cluster is much greater than the lifetime the magnetar. Consequently, by measuring the age of the star-cluster we can determine the age of the progenitor star when it went SN. Then, as a star's lifetime is a strong function of its initial mass, we can estimate the initial mass of the magnetar's progenitor. In the cases of the magnetars SGR~1806-20 and CXOU~J164710.2-455216, associated with the clusters Cl~$1806-20$ and Westerlund~1 (Wd~1) respectively, it appears that the magnetar progenitors had initial masses $\\ga$40\\msun\\ \\citep{Figer05sgr,Bibby08,Muno06}. These results are therefore consistent with the hypothesis that magnetars decend from the most massive stars. Further supporting evidence for this hypothesis comes from a study of the expanding H{\\sc i} shell around the magnetar 1E~1048.1-5937. When the shell was interpreted as a stellar wind bubble blown by the progenitor, a progenitor mass of 30-40\\msun\\ was inferred \\citep{Gaensler05}.  There is a fourth magnetar, SGR~$1900+14$, which can be used to test this hypothesis. It too is thought to belong to a cluster, which was first recognized for its two bright Red Supergiant (RSG) members \\citep{Vrba96,Vrba00}. However, up until now this cluster has been poorly studied. The only current distance estimate is a spectrophotometric distance which assumes an intrinsic brightness for the RSGs, and so cannot be used to derive accurate luminosities for the RSGs themselves. Consequently, any estimate for the magnetar's progenitor mass based on the RSG luminosities is unreliable. The best evidence for the association of the magnetar and star-cluster comes from the detection of an infrared ring around the source, analysis of which placed the magnetar at the same distance from Earth as the cluster's spectrophotometric distance \\citep{Wachter08}.  In this paper we present a spectroscopic and photometric analysis of the stellar content of the cluster Cl~1900+14. We determine a kinematic distance to the cluster and derive the cluster's age. Ultimately, we are able to establish an estimate for the initial mass of the progenitor of SGR~1900+14. We begin in Sect.\\ \\ref{sec:obs} with a description of the observations and data reduction procedure. In Sect.\\ \\ref{sec:res} we present the results of our analysis, and in Sect.\\ \\ref{sec:disc} we discuss the implication for the evolution of massive stars. ",
        "conclusions": "We present an imaging and spectroscopic study of the host cluster of the magnetar SGR1900+14, with the purpose of deriving the initial mass of the magnetar's progenitor. From analysis of the two bright Red Supergiants in the cluster we determine a kinematic distance of 12.5$\\pm$1.7kpc and extinction of $A_{\\rm V} =12.9\\pm0.5$, which is in good agreement with that previously derived for the magnetar. From the luminosities of the RSGs we determine an age for the cluster of 14$\\pm$1 Myr, and a study of the fainter stellar population of the cluster reveals that it is consistent with being coeval to within the errors. Assuming that the SN that created the magnetar occured very recently in comparison to the age of the cluster, we derive an initial mass of the magnetar's progenitor of $M_{\\rm prog} = 17 \\pm 2$\\msun. We find this estimate is insensitive to parameters such as metallicity and the type of evolutionary models used. This result is significantly lower than for other magnetars with progenitor mass estimates, and challenges the hypothesis that very high initial masses ($\\ga40$\\msun) are required to produce magnetars. Instead, we suggest that some other initial parameter, such as magnetic field strength at birth or the presence of a close binary companion, may be the dominant factor in producing a super-strong magnetic field."
    },
    "0910/0910.4542_arXiv.txt": {
        "abstract": "It is shown, that highly accurate estimation of muon shower content can be performed on the basis of knowledge of only vertical depth of shower maximum \\xmaxv\\ and total signal in ground detector. The estimate is almost independent on primary energy and particle type and on zenith angle. The study is performed for 21500 showers, generated with CORSIKA~6.204 from spectrum $E^{-1}$ in the energy range $\\log10(E)$ [eV]=18.5--20 and uniformly in $\\cos^2{\\theta}$ in zenith angle interval $\\theta=0^\\circ-65^\\circ$ for QGSJET~II/Fluka interaction models. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.3538_arXiv.txt": {
        "abstract": "This paper gives an overview of the attempts to determine the distribution of dark matter in low surface brightness disk and gas-rich dwarf galaxies, both through observations and computer simulations.  Observations seem to indicate an approximately constant dark matter density in the inner parts of galaxies, while cosmological computer simulations indicate a steep power-law-like behaviour. This difference has become known as the ``core/cusp problem'', and remains one of the unsolved problems in small-scale cosmology. ",
        "introduction": "Dark matter is one of the main ingredients of the universe. Early optical measurements of the rotation of spiral galaxies indicated the possible presence of large amounts of dark matter in their outer parts (e.g.\\ \\citealt{Rubin:1978p737}), though in many cases the rotation curve could also be explained by the stars alone (e.g.\\ \\citealt{Kalnajs:1983p87,Kent:1986p1363}). Observations at even larger distances from the galaxy centers, using the 21-cm line of neutral hydrogen, definitively confirmed the mass discrepancy \\citep{Bosma:1978p712,Bosma:1981p704,Bosma:1981p703}.  For an extensive review see \\citet{Sofue:2001p714}.  Most of these early observations concentrated on late-type disk galaxies, which all share the property of having an almost constant rotation velocity in their outer parts (the so-called ``flat rotation curve''). As the dynamical contribution of the stars and the gas is insufficient to explain the high rotation velocities in the outer parts, this implies that most of the observed rotation there must be due to some other material, the ``dark matter''. The observed constant velocity suggests that the dark matter in the outer parts of galaxies has a mass density profile closely resembling that of an isothermal sphere, i.e., $\\rho \\sim r^{-2}$.  In the inner parts of galaxies stars are obviously present and they must be the cause of a (possibly large) fraction of the observed rotation velocity. This therefore leads to a transition from the inner parts where the stars contribute to (and in many cases dominate) the dynamics, to the outer parts where the dark matter is important (e.g., \\citealt{Kalnajs:1983p87,Kent:1986p1363,Kent:1987p1410,vanAlbada:1986p701}).   The rotation velocity associated with dark matter in the inner parts of disk galaxies is found to rise approximately linearly with radius. This solid-body behaviour can be interpreted as indicating the presence of a central core in the dark matter distribution, spanning a significant fraction of the optical disk. Some authors adopt a non-singular isothermal sphere to describe this kind of dark matter mass distribution (e.g., \\citealt{Athanassoula:1987p1446}), while others prefer a pseudo-isothermal sphere (e.g., \\citealt{Begeman:1991p1454,Broeils:1992p1452}).  Both models describe the data well (see \\citealt{Kormendy:2004p1455}), and by the late 1980s they had become the \\emph{de facto} description of the distribution of dark matter in (gas-rich, late-type) dwarfs and disk galaxies.  In this paper we use the \\emph{pseudo-isothermal} (PI) sphere to represent the cored models, though this particular choice does not affect any of the discussion in this paper.  The mass density distribution of the PI sphere is given by: \\begin{equation} \\rho_{\\rm PI}(r) = {{\\rho_0}\\over{1+(r/R_C)^2}}, \\end{equation} where $\\rho_o$ is the central density, and $R_C$ is the core radius of the halo. This density distribution leads to an asymptotic flat velocity $V_{\\infty}$ given by $V_{\\infty} = (4\\pi G \\rho_0 R_C^2)^{1/2}$, where $G$ is the gravitational constant.  In the early 1990s, the first results of numerical $N$-body simulations of dark matter halos based on the collisionless cold dark matter (CDM) prescription became available. These did not show the observed core-like behaviour in their inner parts, but were better described by a steep power-law mass-density distribution, the so-called \\emph{cusp}. Fits to the mass-distributions as derived from these early simulations \\citep{Dubinski:1991p4,Navarro:1996p479, Navarro:1997p482} indicated an inner distribution $\\rho \\sim r^{\\alpha}$ with $\\alpha=-1$. (In the following we will use $\\alpha$ to indicate the inner mass density power-law slope.)  The results from these and later simulations are based on the $(\\Lambda)$CDM paradigm, where most of the mass-energy of our universe consists of collisionless CDM in combination with a cosmological constant $\\Lambda$.  This $\\Lambda$CDM paradigm provides a comprehensive description of the universe at large scales (as shown most recently by the Wilkinson Microwave Anisotropy Probe (WMAP) results; see \\citealt{Spergel:2007p773}). However, despite these great successes, it should be kept in mind that the cusp and the central dark matter distribution are not predicted from first principles by $\\Lambda$CDM. Rather, these properties are derived from analytical fits made to dark-matter-only numerical simulations. While the quality and quantity of these simulations has improved by orders of magnitude over the years, there is as yet no ``cosmological theory'' that explains and correctly predicts the distribution of dark matter in galaxies from first principles.   The value $\\alpha = -1$ found in the early CDM simulations is very different from that expected in the PI model, where the constant density core ($\\rho \\sim r^0$) implies $\\alpha = 0$.  These two cases thus lead to two very different descriptions of the dark matter distribution in galaxies. The ``cusp'' ($\\alpha=-1$) models gives rise to a rapidly increasing ``spiky'' dark matter density towards the center, while the ``core'' model ($\\alpha = 0$) has an approximately constant dark matter density. The cusp model therefore has a rotation curve that will rise as the square-root of the radius, while the core model rotation curve rises in a linear fashion. The difference in shapes between the rotation curves of both models is quite pronounced, and, in principle, it should therefore be possible to identify CDM haloes in real galaxies by measuring their rotation curves.  Over the last 15 years or so, much effort has been put into determining the central mass distribution in galaxies using their rotation curves, and comparing them with the outcomes of ever more sophisticated numerical simulations. To first order, one can summarize this work as observational determinations yielding slopes $\\alpha \\sim 0$, while simulations produce $\\alpha \\sim -1$ slopes. This persistent difference is known as the ``core/cusp controversy'', sometimes also described as ``the small-scale crisis in cosmology''. The attempts to reconcile the observations and simulations, either by trying to improve them, or by trying to quantify systematic effects or missing physics, are the subjects of this paper.  I give a brief overview of past and present work dealing with the determination of the central dark matter density distribution in galaxies, with an emphasis on the observational efforts. An overview like this, touching on many different topics in galaxy evolution, cosmology and computational astrophysics, is never complete, and only a small (but hopefully somewhat representative) fraction of the many papers relevant to this topic can be referred to in the limited space available. The rest of this paper is organised as follows: Sect.~2 gives a description of the results that numerical simulations have produced over the years. Section 3 deals with the observational determinations of the dark matter density distribution. Section 4 discusses physical scenarios that have been proposed to reconcile the core and cusp distributions. Section 5 briefly summarizes the work discussed. ",
        "conclusions": "Rotation curves of LSB and late-type, gas-rich dwarf galaxies indicate the presence of constant-density or mildy cuspy ($\\alpha \\sim -0.2$) dark matter cores, contradicting the predictions of cosmological simulations. The most recent simulations still indicate resolved mass density slopes that are too steep to be easily reconciled with the observations (typically $\\alpha \\sim -0.8$ at a radius $\\sim 0.1$ kpc).  Claims of shallow slopes at even smaller radii depend on the validity of the analytical description chosen for the mass-density profile.  Whereas early HI observations and long-slit H$\\alpha$ rotation curves still left some room for observational (pointing, resolution) or physical (non-circular motions, tri-axiality) systematic effects to create the illusion of cores in the presence of a cuspy mass distribution, the high-resolution optical and HI velocity fields that have since become available significantly reduce the potential impact of these effects. Measured non-circular motions and potential ellipticities are too small to create the illusion of a core in an intrinsically cuspy halo.  This indicates either that halos did not have cusps to begin with, or that an as yet not understood subtle interplay between dark matter and baryons wipes out the cusp, where the quiescent evolution of LSB galaxies severely limits the form this interplay can take. Adiabatic contraction and dynamical friction yield contradictory results, while models of massless disks in tri-axial halos result in preferred viewing directions.  LSB galaxy disks, despite their low $\\Upsilon_{\\star}$ values, are not entirely massless, and observations and simulations will need to take this into account. Similarly, the difficulties in reconciling a possible underlying triaxial potential with the circularizing effects of the baryons also needs to be investigated. In short, studies which, constrained and informed by the high-quality observations now available, self-consistently describe and model the interactions between the dark matter and the baryons in a cosmological context are likely the way forward in resolving the core/cusp problem."
    },
    "0910/0910.3823_arXiv.txt": {
        "abstract": "The Extreme-Ultraviolet Imaging Spectrometer on board the HINODE satellite is used to examine the loop system described in \\cite{me09} by applying spectroscopic diagnostic methods. A simple isothermal mapping algorithm is applied to determine where the assumption of isothermal plasma may be valid, and the emission measure locii technique is used to determine the temperature profile along the base of the loop system. It is found that, along the base, the loop has a uniform temperature profile with a mean temperature of $0.89 \\pm 0.09$~MK which is in agreement with the temperature determined seismologically in \\cite{me09}, using observations interpreted as the slow magnetoacoustic mode. The results further strengthen the slow mode interpretation, propagation at a uniform sound speed, and the analysis method applied in \\cite{me09}. It is found that it is not possible to discriminate between the slow mode phase speed and the sound speed within the precision of the present observations. ",
        "introduction": "The temperature of the plasma within coronal loops has traditionally been determined using spectroscopic methods such as emission line ratios \\citep{phi08} and emission measure locii techniques \\citep{jor87, lan02, del02}. In \\cite{me09}, the temperature along a coronal loop structure was seismologically determined using Solar Terrestrial Relations Observatory/Extreme-Ultraviolet Imager \\citep[STEREO/EUVI,][]{wue04} observations of wave propagation along a loop system. These waves were interpreted as manifestations of the slow magnetoacoustic mode. The stereoscopic observations were used to derive the propagation geometry with an inclination of ${37 \\pm 6} ^{\\circ}$ to the local normal and true coronal slow mode phase speed of $132 \\pm 9$ km~s$^{-1}$. Thus, the sound speed was then used to infer the plasma temperature of $0.84 \\pm 0.12$~MK.  \\cite{me09} was the first direct measurement of the slow mode speed within a coronal loop and inference of the loop plasma temperature using this technique. The work presented here aims to provide an independent observational test of those results and conclusions. The plasma temperature measured using spectroscopic emission line diagnostics is compared to the result obtained using the seismological technique, and confirmed to be in agreement. ",
        "conclusions": "The results presented here determine, spectroscopically, the temperature of the active region loop system presented in \\cite{me09}, along which slow magnetoacoustic waves were found to propagate. The seismological technique, applied in \\cite{me09}, derived a plasma temperature of  $0.84 \\pm 0.12$~MK. This temperature is independently confirmed, here, using the emission measure locii technique, with a derived temperature of $T =0.89 \\pm 0.09$~MK, and is consistent with the \\cite{me09} results.  The agreement between the results validates the technique applied in \\cite{me09}, and further strengthens the slow magnetoacoustic mode interpretation of the observed waves. The consistency between the two independent estimates of the temperature also suggests that the assumption of an isothermal plasma in the loop system is valid. The temperature measured at the base of the loop system shown in Figure~\\ref{fig5} displays a uniform temperature profile. This result is also in agreement with the EUVI observations presented in \\cite{me09}, where the observed waves have a constant phase speed as a function of distance. Thus, indicating a uniform temperature profile along the base of the loop system, at least to the extent of where the waves are observed.  \\cite{wan09} recently investigated waves similar to those reported in \\cite{me09} and found consistent results. Using HINODE/EIS observations interpreted as slow magnetoacoustic waves at the footpoint of a coronal loop, they use the Doppler shift amplitude to estimate the loop inclination and estimate a temperature of $0.7 \\pm 0.3$~MK. Again, this is consistent with the results of the spectroscopic diagnostics applied here.  In \\cite{me09}, it was suggested that it may be possible to use the propagating slow mode to measure the coronal magnetic field strength. The HINODE/EIS observations presented here indicate that, at least within the current observational diagnostic precision, at the temperature and density of this region, it is not possible to discriminate the slow mode tube speed from the sound speed to make such a measurement. This may be possible in higher temperature structures where, due to its dependence on $T$, the sound speed would be expected to have a greater divergence from the tube speed. This may be a possible achievable goal of the Solar Orbiter mission, with high resolution observations of wave propagation in different structures."
    },
    "0910/0910.2708_arXiv.txt": {
        "abstract": "The gravitational potential of clusters of galaxies acts as a cosmic telescope allowing us to find and study galaxies at fainter limits than otherwise possible and thus probe closer to the epoch of formation of the first galaxies.  We use the Bullet Cluster {\\bullet} ($z = 0.296$) as a case study, because its high mass and merging configuration makes it one of the most efficient cosmic telescopes we know.  We develop a new algorithm to reconstruct the gravitational potential of the Bullet Cluster\\, based on a non-uniform adaptive grid, combining strong and weak gravitational lensing data derived from deep \\HST/ACS F606W-F775W-F850LP and ground-based imaging.  We exploit this improved mass map to study $z\\sim 5-6$ Lyman Break Galaxies (LBGs), which we detect as dropouts. One of the LBGs is multiply imaged, providing a geometric confirmation of its high redshift, and is used to further improve our mass model. We quantify the uncertainties in the magnification map reconstruction in the intrinsic source luminosity, and in the volume surveyed, and show that they are negligible compared to sample variance when determining the luminosity function of high-redshift galaxies.  With shallower and comparable magnitude limits to HUDF and GOODS, the Bullet cluster observations, after correcting for magnification,  probe deeper into the luminosity function of the high redshift galaxies than GOODS and only slightly shallower than HUDF. We conclude that accurately focused cosmic telescopes are the most efficient way to sample the bright end of the luminosity function of high redshift galaxies and - in case they are multiply imaged - confirm their redshifts. ",
        "introduction": "\\label{sec:intro}  Studying galaxies at high redshifts is crucial for understanding both their formation and evolution, and the role they played in cosmic reionization.  However, recent observations of $z\\gtrsim 6$ objects (e.g., \\citealp{bouwens08}) show that UV-bright galaxies alone are not believed to be sufficient to reionize the Universe (unless for example, the escape fractions or clumping factors are dramatically different from current assumptions). It is therefore important to continue to find radiation sources at the highest redshifts and study their physical properties.   Observations of galaxies at these high redshifts are challenging, not only due to the large luminosity distance to these objects, but also due to their lower luminosity, since the luminosity function evolves compared to galaxies at moderate redshifts ($z\\sim2$; see e.g. \\citealp{reddy09,bouwens07}).  Recently our knowledge of the early cosmic epoch of galaxy formation ($z\\gtrsim 4$) has vastly increased, providing an opportunity for us to test theories of star formation \\citep[and references therein]{stark09,bouwens08, bouwens07, yoshida06, stanway05, giavalisco04}. Comparing the measured luminosity function with results from simulations provides insight on the efficiency of star formation as a function of mass scale, and on the feedback mechanisms that may limit it (see e.g. \\citealp{vandenbosch03}).   We can find high-redshift galaxies by searching for the redshifted Lyman break using broad band photometry \\citep{steidel96}.  These sources, known as Lyman Break Galaxies (LBGs, see e.g. \\citealp{vanzella09} for recent studies, and \\citealp{giavalisco02} for a review) are the best studied and the largest sample of galaxies at redshifts $z\\gtrsim 5$.  One can identify $z\\simeq 5$ objects by their non-detection in the $V$-band and blueward (e.g. \\citealp{giavalisco04b,stanway05,hildebrandt09, stark09}): such objects are referred to as ``$V$-band dropouts.'' Similarly, objects at $z\\simeq 6$ are associated with $i$-band non-detection (e.g. \\citealp{stanway03, bouwens08}), and $z\\simeq 7$\\nobreakdash--$8$ with $z$-band (e.g. \\citealp{bouwens06,bouwens08, henry09, oesch09}). To date, preliminary results using the new WFC3/IR data have now been released. \\citet{oesch09b} and \\citet{bunker09} also study the $z$-band dropouts, whereas \\citet{bouwens09b}, \\citet{yan09}, and \\citet{mclure09} study $Y$-band dropouts which are potential $z\\sim 8-8.5$ objects. At even higher redshifts, $z\\gtrsim 8$ objects are observed as $J$-band dropouts (e.g. \\citealp{yan09}, \\citealp{henry08}, and \\citealp{bouwens05}).\\footnote{The $z\\sim9$ candidate object from \\protect\\citealp{henry08} was later confirmed to be a lower redshift one and the number of candidates from \\citealp{bouwens05} appears consistent with the expected contamination from low-redshift interlopers.} The main limitations of these experiments to date are the relatively low expected number counts (even in the deepest fields), and contamination ($z \\sim 2$ dusty galaxies and L and T dwarf stars can also share the characteristic dropout appearance). Furthermore, the depths probed by these surveys are around characteristic luminosity for high-redshift population $L_{*}$ and do not reach far into the low luminosity regime where the faint-end slope is constrained, especially at higher redshifts.   We argue in this paper that at the luminosity regime down to $L\\sim L_{*}$, determining the high redshift luminosity function can be performed more efficiently when using galaxy clusters as cosmic telescopes with respect to blank fields. This technique was proposed shortly after the first gravitationally-lensed arcs in galaxy clusters were discovered \\citep{soucail90} and has been successfully applied to ever-improving data (e.g., \\citealp{zheng09,bouwens09,richard06,ellis01}). Observations of galaxy clusters used as cosmic telescope is consistently delivering record holders in the search for the highest redshift galaxies \\citep{kneib04,bradley08}. The reason for this is that the magnifying power of a massive cluster lens typically allows the detection of objects more than a magnitude fainter than the observation limit. Due to this same magnification, the effective solid angle of the survey volume decreases; however, since the luminosity function is practically exponential at the magnitudes we probe, the lensing magnification gives us a substantial net gain. Indeed, if the faint end slope of the luminosity function at high redshift is steep (e.g. \\citealp{bouwens07}), we gain at magnitudes below and around the characteristic magnitude $M^*$. One concern is that we need to properly account for the errors introduced when we determine the true observed volume with a cosmic telescope (e.g. \\citealp{bouwens09}). However, we quantify the uncertainties in the magnification map and show that they are negligible compared to sample variance when determining the luminosity function of high-redshift galaxies.   Cosmic telescopes offer two further advantages.  First, the detected sources behind a cluster lens are magnified in apparent size. As a result, we can resolve smaller physical scales than would otherwise be possible, and begin to actually measure the properties of $z\\gtrsim6$ galaxies on a case-by-case basis. Resolving the galaxies also helps in rejecting contamination by foreground cold stars. Second, some of the sources will be multiply imaged. If this is the case, then they can be readily distinguished from the main contaminants, typically intermediate redshift galaxies and cold stars. The positions, where multiply imaged systems form are redshift dependent, so a $z>5$ multiply imaged source can be readily distinguished from (multiply-imaged) $z\\sim 2$ source and singly imaged stars.  Thus, in the magnitude range we observe here, a cosmic telescope survey is more efficient than a competing imaging survey in unlensed fields to the same depth.   The most effective gravitational lenses for such studies are massive and highly elongated (in projection, on the plane of the sky) galaxy clusters. It is easy to understand the first requirement, as more massive objects are more powerful lenses. The ellipticity is important as well, as the area of highest magnification is located at the semi major-axis intersection with the critical curve (i.e. curve connecting points of formally infinite magnification).  Equivalently, sources close to the major axis cusp have the largest magnification. The cluster of galaxies {\\bullet}, discovered by \\citet{tucker95}, is one of the hottest, most X-ray luminous clusters known. It is also a plane-of-the sky merger (e.g. \\citealp{markevitch02,barrena02}), hence giving a highly elongated mass distribution in projection.   It is also at a favorable redshift, which contributes to make it a very efficient lens. The critical curve radius decreases with increasing lens redshift for a source at a fixed (higher) redshift. One might thus think that lower redshift clusters $0.05< z \\lesssim 0.2$ are better suited for these surveys. However, a small solid angle coverage by cluster member galaxies is desirable, hence reducing the effectiveness of clusters that are at too low redshift. Whereas lensing efficiency scales with angular diameter distances between observer, source, and the lens, the size of cluster members only depends upon the latter. There is an optimal balance between the size of the critical curve and the size of the cluster members. In addition, a good cosmic telescope candidate should also show many strongly lensed images (i.e. efficient also for lensing of $z\\sim 1-3$ sources), which are needed to reliably reconstruct its magnification map. Bullet Cluster, with $>10$ definite strongly lensed systems and the lens redshift of $z\\sim 0.3$,  is well suited for these studies.  Specifically, since we are trying to detect objects close to the detection threshold, we do not want to rely on perfect image subtraction of the cluster members.  For clusters at much higher redshifts, on the other hand, the critical curves (and the area of high magnification) are not large enough.   A key requirement for a cosmic telescope survey is an absolutely calibrated, high-resolution mass reconstruction for the lens. This is necessary to accurately convert the observed solid angle into the actual cosmic volume, as well as to reconstruct the source fluxes and so determine their luminosity. In this paper we will describe a newly developed technique to reconstruct the mass distributions of galaxy clusters.  As we will show, the technique allows us to improve the accuracy in the very center of this cluster, which is crucial to accurately predict the magnifications of dropouts (and other objects in the field). The key feature of this new technique is the optimal use of information from both multiple image systems (strong lensing from deep high-resolution \\HST/ACS data) and from the distortions of the singly-imaged background sources (weak lensing, from combined ACS and ground based data) using adaptive grid. Multi-resolution methods have been used in either strong or weak gravitational lensing in the past (see e.g. \\citealp{marshall06,diego07,merten08,deb08,jullo09}). The new method described here, uses information from both strong and weak lensing regimes, on a multi-level grid that allows us to efficiently reconstruct features at the maximum resolution allowed by the data. This new method and map supersedes our previously-inferred potential map, which was limited in its resolution due to the inability of the mass-modeling routine to simultaneously reconstruct large angular scale features (constrained by the relatively low signal-to-noise weak lensing data), and the smaller-scale features in the high signal-to-noise strong lensing regime.    This paper is structured as follows. Section~\\ref{sec:swunitedamr} describes the improved method to reconstruct the mass distribution on a non-uniform grid of pixels. Section~\\ref{sec:data} describes the data and its reduction procedures, and in Section~\\ref{sec:results} we present the mass reconstruction of Bullet Cluster. Section~\\ref{sec:dropouts} then gives the results of our search for $z\\sim 5-6$ objects using our newly-calibrated lens, and discusses further the advantages and disadvantages of performing such studies behind clusters.  Our conclusions are summarized in Section~\\ref{sec:conclusions}.  Throughout the paper we assume a $\\Lambda$CDM cosmology with $\\Omega_{\\rm m} = 0.3$, $\\Omega_{\\Lambda} = 0.7$, and Hubble constant $H_0 = 70 {\\rm \\ km \\: s^{-1}\\:\\mbox{Mpc}^{-1}}$. At the cluster redshift $z_{\\rm cl}=0.296$, the physical length scale is 4.4 kpc/arcsec. All the coordinates in this paper are given for the epoch J2000.0. All magnitudes are given in the AB system. ",
        "conclusions": "\\label{sec:conclusions} We have presented a new mass reconstruction for the post-collision Bullet Cluster. In addition, we have used its extraordinary capabilities as a cosmic telescope to study $z\\sim 5-6$ objects. Our main conclusions are as follows: \\begin{enumerate} \\item Our new strong and weak lensing mass reconstruction method, that uses a non-uniform adaptive pixel grid, performs very well in reconstructing the cluster potential in the vicinity of multiple images. In particular, we have improved the accuracy with which the image positions are reproduced to an average rms residual of $1.4\\arcsec$. The limitation in pushing this number to the observed uncertainties in image position ($\\sim 0.1\\arcsec$) is mostly in the lack of secure redshifts (only one spectroscopic redshift is known for the 12 multiple image systems used here). The method can resolve the small scale substructure and reconstruct the images even more precisely, once the redshifts become available. \\item Bullet Cluster is an efficient cosmic telescope. We have used its power to search for $z\\sim 5-6$ objects using the dropout technique. After correcting for the lensing effect and taking into account the errors from the imperfect magnification map, we derive the surface densities of $V$ and $i$-band dropouts. They match well the densities (in absolute value) derived from previous studies \\citep{stark09,bouwens07,beckwith06} using either deeper or larger surveys (HUDF and GOODS). E.g. GOODS (release 1) data has similar exposure times ($5000\\mbox{s}$ in $V$ and $i$ and $10660$ in $z$) over $360\\mbox{ arcmin}^2$ field. Scaling total numbers of dropouts from \\citet{stark09} (using area ratio between our and their survey) we would obtain 14 $V$ and 3.5 $i$-band dropouts; whereas we observe 20 and 4. In addition they are at fainter (intrinsic) magnitudes as their sample. \\item Our sample numbers are much smaller (as our search area is $10\\mbox{ arcmin}^2$, compared to the area used in e.g. \\citet{stark09,bouwens07} with $>300\\mbox{ arcmin}^2$), and so the errorbars are dominated by sample variance and Poisson noise. However, our results show, that at the same observed magnitudes, we gain in the number of sources observed compared to a blank field of the same size. In addition, we observe intrinsically fainter sources than would otherwise be possible in a blank field observed to similar depths. Finally, multiply imaged sources are easily discriminated from contaminants, as lens model roughly predicts their redshifts, and hence can be distinguished from $z\\sim2$ dusty galaxy and cool stars. \\end{enumerate}  This analysis can and will be extended to higher redshift objects (z and J-dropouts).  The main requirement for such studies is deep optical and near-IR data, as well as high resolution magnification maps such as the one presented in this paper.  These studies can also be readily expanded to a larger sample of clusters. With a sample of $\\sim 20$ massive clusters, with well understood magnification properties and observed to similar depths as the Bullet Cluster we can increase the expected number counts and reduce the errors due to sample variance and Poisson sampling to those of GOODS (see e.g. \\citealp{stark09}). However, due to lensing, we can push observations to at least a magnitude deeper than GOODS and  only slightly shallower than HUDF (in much shorter observing time). Finally, highly magnified images ($\\mu\\gtrsim 50$) potentially allow for (otherwise prohibitive) spectroscopic follow up of these high redshift galaxies.  Highly magnified objects will often sample populations that would otherwise be unobservable by any practical means."
    },
    "0910/0910.0011_arXiv.txt": {
        "abstract": "I present a predictive analysis for the behavior of the far-infrared (FIR)--radio correlation as a function of redshift in light of the deep radio continuum surveys which may become possible using the square kilometer array (SKA). To keep a fixed ratio between the FIR and predominantly non-thermal radio continuum emission of a normal star-forming galaxy, whose cosmic-ray (CR) electrons typically lose most of their energy to synchrotron radiation and Inverse Compton (IC) scattering, requires a nearly constant ratio between galaxy magnetic field and radiation field energy densities. While the additional term of IC losses off of the cosmic microwave background (CMB) is negligible in the local Universe, the rapid increase in the strength of the CMB energy density (i.e. $\\sim(1+z)^{4})$ suggests that evolution in the FIR-radio correlation should occur with infrared (IR;~$8-1000~\\micron$)/radio ratios increasing with redshift. This signature should be especially apparent once beyond $z\\sim3$ where the magnetic field of a normal star-forming galaxy must be  $\\sim$50~$\\mu$G to save the FIR-radio correlation. At present, observations do not show such a trend with redshift; $z\\sim6$ radio-quiet quasars appear to lie on the local FIR-radio correlation while a sample of $z\\sim4.4$ and $z\\sim2.2$ submillimeter galaxies (SMGs) exhibit ratios that are a factor of $\\sim$2.5 {\\it below} the canonical value. I also derive a 5$\\sigma$ point-source sensitivity goal of $\\approx$20~nJy (i.e. $\\sigma_{\\rm RMS} \\sim 4$~nJy) requiring that the SKA specified be $A_{\\rm eff}/T_{\\rm sys}\\approx  15000$~m$^{2}$~K$^{-1}$; achieving this sensitivity should enable the detection of galaxies forming stars at a rate of $\\ga25~M_{\\sun}~{\\rm yr}^{-1}$, such as typical luminous infrared galaxies (i.e. $L_{\\rm IR} \\ga 10^{11}~L_{\\sun}$), at all redshifts if present. By taking advantage of the fact that the non-thermal component of a galaxy's radio continuum emission will be quickly suppressed by IC losses off of the CMB, leaving only the thermal (free-free) component, I argue that deep radio continuum surveys at frequencies $\\ga$10~GHz may prove to be the best probe for characterizing the high-$z$ star formation history of the Universe unbiased by dust. ",
        "introduction": "Radio continuum emission from galaxies arises due to a combination of thermal and non-thermal processes primarily associated with the birth and death of young massive stars, respectively. The thermal (free-free) radiation of a star-forming galaxy is emitted from H{\\sc ii} regions and is directly proportional to the photoionization rate of young massive stars. Since emission at GHz frequencies is optically thin, the thermal radio continuum emission from galaxies is a very good diagnostic of a galaxy's massive star formation rate.  Massive ($\\ga 8~M_{\\sun}$) stars which dominate the Lyman continuum luminosity also end their lives as supernovae (SNe) whose remnants (SNRs) are responsible for the acceleration of cosmic-ray (CR) electrons into a galaxy's general magnetic field resulting in diffuse synchrotron emission. Thus, in a more complicated manner, the non-thermal radio continuum emission also traces the most recent star formation activity in a galaxy. Ideally, one would like to isolate the thermal component as it is a more direct measure of the most  recent massive star formation activity. However, at GHz frequencies, the non-thermal fraction typically dominates the total radio continuum emission \\citep[i.e. $\\sim$10:1 at $\\sim$1~GHz;][]{cy90} making the isolation of the thermal fraction difficult.    These same massive stars are often the primary sources of dust heating in the interstellar medium (ISM) as their starlight is absorbed and reradiated at far-infrared (FIR) wavelengths by interstellar grains. This common origin between the FIR dust emission and thermal $+$ non-thermal radio continuum emission from galaxies is thought to be the dominant physical processes driving the FIR-radio correlation on global \\citep[e.g.][]{de85,gxh85,sn97,nb97,yrc01} and local \\citep[e.g.][]{bg88,xu92,mh95,hoer98,hip03,ejm06,ejm08,ah06} scales. While massive star formation provides a shared origin between emission at these two wavelengths, a number of physical processes must conspire to yield such a remarkably constant ratio among galaxies spanning nearly 5 orders of magnitude in luminosity  \\citep[e.g.][]{yrc01}, among other properties (e.g. Hubble type, FIR color, and FIR/optical ratio). For example, since the non-thermal synchrotron emissivity of a single electron is roughly proportional to the electron density times the square of the magnetic field strength, it would seem that the FIR-radio correlation should be highly dependent on the propagation of CR electrons and a galaxy's magnetic field distribution and strength. If magnetic fields in galaxies are built-up over time, and therefore were weaker in the past, the correlation should be different at higher redshifts. However, all indications suggest that the FIR-radio correlation holds out to moderate redshifts \\citep[e.g.][]{mg02,grup03,pa04,df06,ejm09a,ms09} suggesting that the magnetic field strength and structure in these galaxies is similar to what is observed for galaxies in the local Universe. Consequently, comparing the FIR and radio emission characteristics of galaxies can illuminate properties that are typically inaccessible by other observational means.  \\begin{figure*} \\plottwo{f1a.eps}{f1b.eps} \\caption{ {\\it Left:} The cooling timescales of CR electrons as a function of electron energy due to a number of physical processes: synchrotron, Inverse Compton (IC), ionization, and bremsstrahlung. For these calculations typical parameters for normal star-forming galaxies are assumed: $B=10~\\mu$G, $U_{\\rm rad} = 10^{-12}$~erg~cm$^{-3}$, and $n_{\\rm ISM} = 1$~cm$^{-3}$. A timescale due to the escape of CR electrons from a galaxy, assuming random walk diffusion and an escape scale-length of $l_{\\rm esc} = 3$~kpc, is also included. The top axis indicates the synchrotron emitting frequency of the CR electrons. The shaded region indicates the electron energies emitting predominantly in the $0.5-2$~GHz passband. {\\it Right:} Same as what is plotted in the left panel, but for values that may be more typical of galaxies hosting compact starbursts, namely: $B=1$~mG and $n_{\\rm ISM} = 10^{4}$~cm$^{-3}$, which assumes the usual $B \\propto \\sqrt{n_{\\rm ISM}}$ scaling. The additional term from adiabatic expansion losses due to a starburst-driven galactic wind, having a velocity of $v_{w} = 300$~km~s$^{-1}$ and assuming a disk scale-height $h= 1$~kpc, are also shown. $U_{\\rm rad}$ was also scaled up to 10$^{-8}$~erg~cm$^{-3}$. For normal star-forming galaxies, synchrotron and IC losses dominate CR electron energy losses, along with escape, for radiation observed at $\\ga$GHz frequencies. However, in the case where a mG magnetic field is considered, synchrotron and IC losses no longer are the dominant cooling processes for CR electrons emitting at GHz frequencies. \\label{fig-1}} \\end{figure*}    At present, detailed studies of the FIR-radio correlation among star-forming galaxies has been limited to only moderate redshifts at $z \\la 1$. With the sensitivity of existing FIR and radio capabilities, studies at higher redshifts only probe the most extreme objects. For example, deep ($\\sigma_{\\rm RMS} \\sim0.55~$mJy) 70$\\micron$ observations \\citep{df06} from the the Far-Infrared Deep Extragalactic Legacy (FIDEL; PI. M. Dickinson) survey is able to probe luminous infrared galaxies (LIRGs; $10^{11} \\leq L_{\\rm IR} < 10^{12}~L_{\\sun}$) and ultraluminous infrared galaxies  (ULIRGs;  $\\geq 10^{12}~L_{\\sun}$) out to redshifts of $z\\sim 1$ and $z\\sim2 $, respectively. With {\\it Herschel}, surveys such as GOODS-{\\it Herschel} should improve this situation by measuring the FIR properties of normal star-forming galaxies out to $z\\sim 1$, and that of LIRGs and ULIRGs out to redshifts of $z\\sim2$ and $z\\sim4$, respectively. ALMA will also play a significant role by measuring the peak of the FIR spectral energy distribution (SED)  for all LIRGs at $z\\ga5$.  These populations of dusty star-forming galaxies appear to dominate the stellar mass assembly at increasing redshifts; the star formation rate density increases by a factor of $\\sim5-10$ between $z\\sim0$ and $z\\sim1$,  becoming increasingly obscured (i.e. $\\ga60$\\%) by dust \\citep{de02,ce01,el05,bm09}. Thus, LIRGs and ULIRGs appear to dominate the luminosity density at increasing redshifts, and optically thin measures of star formation and active galactic nuclei (AGN) activity are critical for proper quantification of stellar mass build-up over cosmic time. At radio wavelengths, however, detecting such high redshift galaxies remains extremely difficult. Even with a fully operational EVLA, IR-bright star-forming galaxies  (e.g. M~82; $L_{\\rm IR} \\approx 4\\times10^{10}~L_{\\sun}$) and moderate LIRGs  (i.e. $L_{\\rm IR} \\approx 3\\times10^{11}~L_{\\sun}$) will not be detectable beyond redshifts of $z\\sim 1$ and $z\\sim 2$, respectively.  A next-generation radio facility such as the Square Kilometer Array (SKA) should easily remedy this disparity between the depth of FIR and radio continuum surveys. While the SKA was initially proposed solely on the basis of H{\\sc i} science, deep continuum imaging is critical for the realization of nearly all of the five established Key Science Projects (KSPs), especially  studies of ``The Origin and Evolution of Cosmic Magnetism\" and ``Galaxy Evolution and Cosmology.\" Essential to these two science goals is the proper measurement of the star formation and AGN history over cosmic time for which deep radio continuum studies may provide an excellent advantage over other wavelengths.  In this paper I present physically motivated expectations for the behavior of the FIR-radio correlation at increasing redshift and discuss their implications. I also discuss how the SKA, when combined with future FIR observations to be obtained with next generation facilities both prior to, and commensurate with, final SKA science operations, will be able to tackle interesting problems associated with these KSPs. The paper is organized as follows: In $\\S$2 I discuss the FIR-radio correlation and introduce the physical processes for which it depends on. Then, in section $\\S$3, I present the results for the expected evolution in the FIR-radio correlation at increasingly high redshifts and report on the technical specifications for detecting high-$z$ galaxies in deep radio continuum surveys. In Section $\\S$4 I discuss the physical implications of the results for characterizing the properties of galaxies at high~$z$ and compare these theoretical expectations with existing observations. Finally, in $\\S$5, I summarize my conclusions. ",
        "conclusions": "I will now discuss how various galaxy properties can be inferred from discrepancies between expected and observed IR/radio ratios. I will also discuss what this means for the prospect of using deep radio continuum surveys, both at high and low frequencies, to study star formation over cosmic time using the SKA.  \\subsection{Searching for the Emergence of Large-Scale Magnetic Fields} One of the primary energetic constituents of the ISM is its magnetic field. Its energy density is equal to that of a galaxy's radiation field, CRs, and turbulent gas motions, yet its origin and role in shaping galaxy evolution, in particular through the possible regulation or star formation, remains highly uncertain  \\citep[e.g.][]{kf01,cox05}. At present the leading theory for large-scale magnetic field generation is the mean-field dynamo model \\citep[e.g.][]{rb96}. In this paradigm galaxy-scale magnetic fields are built up over time, suggesting that field strengths should be typically much weaker in the past. The exact details on how the amplification of weak, seed magnetic fields in the halos of protogalaxies eventually translates into galaxy-scale fields is uncertain. Does the first epoch of star formation quickly initiate a dynamo in a galaxy? Does the field form during the in-fall of gas prior to star formation? Recent theoretical work suggests that weak, seed magnetic fields may be quickly amplified by the turbulent (small-scale) dynamo, yielding strong ($\\sim$20~$\\mu$G) turbulent fields by $z\\sim 10$ \\citep{ta09}. These fields then serve as a seed for the mean-field (large-scale) dynamo, resulting in regular fields of $\\mu$G strengths having a few kpc coherence within $\\sim$2~Gyr (i.e. by $z\\sim3$). Such a scenario is consistent with indirect evidence from Faraday rotation studies of distant radio sources which claim $\\mu$G fields to be in place for normal galaxies at $z\\sim1.5$ \\citep{pk08,mb08}.   Since the non-thermal radio emission from galaxies, and therefore the linearity of the FIR-radio correlation, depends critically on the magnetic field strength of a galaxy, deviations from nominal ratios as a function of redshift may help to illuminate the presence and characteristics of magnetic fields in galaxies at early epochs. However, given that regular magnetic fields are thought to weaker in the past, as suggested by the mean-field dynamo model, maintaining a constant IR/radio ratio with increasing redshift by increasing magnetic field strength with redshift is somewhat contrived. If the FIR-radio correlation is found to hold for normal galaxies at $z\\ga$3, either the the turbulent dynamo and associated small-scale magnetic fields play an important role at early epochs, or our physical picture for the correlation needs revision. This assumes of course that the existence of an AGN in these galaxies, whose radio emission may fortuitously bring the IR/radio ratio to near or below the local correlation value, can be ruled out. Such situations are discussed below in $\\S$\\ref{sec-data} where I compare observations with these expectations.   In the absence of strong magnetic fields at high-$z$, the IR/radio ratios of galaxies should be large, similar to what is observed for `nascent' starbursts. These systems are rare ($\\la$1\\%) in the local Universe \\citep{hr03}, and are thought to be experiencing the onset of an episode of star formation after a long ($\\approx$10$^{8}$~yr) period of quiescence. Since OB stars will always ionize the nearby gas, generating free-free emission, these galaxies may only be identifiable in the radio by their thermal radio continuum emission.    \\begin{figure*} \\plotone{f6.eps} \\caption{ The same as Figure \\ref{fig-5}, except that I have plotted actual data from a number of studies in the literature. A logarithmic redshift scale is used since there are so few points at $z\\ga3$. Filled in symbols indicate that an object is an SMG, i.e. the entire \\citet{ak06}, \\citet{ev07},\\citet{pc08},  \\citet{ed09a,ed09b}, \\citet{kc09}, \\citet{vs09}, \\citet{ak09}, and \\citet{kk09}  samples, as well as half of the  \\citet{ejm09a} sample. The lower limit ($uppward$-$arrow$) for the \\citet{kk09} data point is due to only having an upper limit at 1.4~GHz. The radio-quiet quasar J114816+525150 is labeled; at a redshift of $z=6.64$ it lies at $\\sim$1/16 of the Universe's current age. Rather than seeing galaxies tend toward higher IR/radio ratios per expectations given that IC losses of the CMB become increasingly important with increasing redshift, many galaxies tend to have IR/radio ratios lower than the local value (i.e the SMGs).  % \\label{fig-6}} \\end{figure*}  \\subsection{Current Observational Results \\label{sec-data}} Using existing facilities a number of deep FIR, submm, and mm surveys have detected galaxies at fairly high redshfits. Sources detected in the FIR and radio from a number of these studies are included in Figure \\ref{fig-6} with the highest redshift sources being the $z\\sim6$ quasars detected by \\citet{cc04} and \\citet{ab06}. In cases where only a FIR measurement is provided, the assumption that ${\\rm IR(8-1000~\\micron) \\approx 2\\times FIR(42-122~\\micron)}$ is made, which is exactly the difference between the average FIR/radio and IR/radio ratios reported by \\citet{yrc01} and \\citet{efb03}, respectively. The average and $\\pm~1\\sigma$ lines for the local IR-radio correlation are drawn \\citep[i.e. $q_{\\rm IR} \\approx 2.64 \\pm 0.26$~dex;][]{efb03}. As in Figure \\ref{fig-5}, the expected locations of a typical star forming galaxy as IC losses of the CMB become increasingly important relative to the internal synchrotron losses of CR electrons for various magnetic field strengths are overplotted. Since so few sources are detected at $z\\ga3$, a logarithmic scaling is used for the redhift axis to better identify any trends.  \\subsubsection{Submillimeter Galaxies} Most of the galaxies out to $z\\sim1$ \\citep[e.g. the 70~$\\micron$ detected sources in; ][]{df06,ejm09b} tend to have IR/radio ratios similar to the local value, which is consistent with the expectations presented here since the magnetic fields in such galaxies need not be excessively large to compete with IC losses off of the CMB. Once beyond $z\\sim1$ there is an increasingly larger number of galaxies having IR/radio ratios below the canonical value, rather than larger per the suggested expectations. The majority of these sources are submillimeter galaxies (SMGs). While typically detected in the hard-band ($2-8$~keV) X-rays,  thus indicating the presence of an AGN, this result is somewhat surprising since SMGs are thought to be primarily powered by star formation \\citep[e.g.][]{da05, ejm09a}.  The samples of \\citet{ak06} and \\citet{ev07} consist entirely of SMGs and the filled triangles indicate the SMGs included in \\citet{ejm09a}. The mean redshift among these three samples of SMGs is $z \\approx2.17\\pm0.65$. \\citet{ak06} report IR/radio ratios for their sample of SMGs to be systematically below the local value. \\citet{ejm09a} also found  that the SMGs in their sample to have IR/radio ratios which are, on average, a factor of $\\sim$3.5 lower than the canonical ratio. The mean IR/radio ratio for the sample of \\citet{ev07} is $\\sim2.27\\pm0.16$~dex, a factor of $\\sim$2 below the local value of 2.64~dex. The average IR/radio ratio among all 38 of these $z\\sim$2.2 SMGs is $q_{\\rm IR} = 2.26\\pm0.41$~dex, a factor of $\\sim$2.4 times smaller that the local value. A number of possible explanations for why these systems depart for the FIR-radio correlation are included in \\citet{ejm09a}, but a definitive conclusion has not been drawn.  Recently, a number of higher redshift SMGs have been identified \\citep[][]{pc08,ed09a,ed09b,kc09,vs09,ak09,kk09}. The average redshift among these 8 objects is $z \\approx4.40\\pm0.33$. Similar to the $z\\sim2.2$ SMGs, the mean IR/radio ratio of these $z\\sim4.4$ SMGs is $q_{\\rm IR} = 2.20 \\pm 0.27$~dex, a factor of $\\sim$2.8 times smaller that the local value. This result is somewhat surprising since the expected increase in the IR/radio ratios between a $z\\sim2.2$ and $z\\sim4.4$ galaxy due to the increase in IC losses to CR electrons from the CMB is a factor of $\\sim$3 for a $\\approx$35~$\\mu$G field, the average value for the \\citet{ejm09a} SMGs using the revised minimum energy formula of \\citet{bk05}.  While it has been suggested that the properties of the $z\\sim4.4$ SMGs are similar to the  $z\\sim2.2$ SMGs \\citep[e.g.][]{ed09a}, the consistency in their IR/radio ratios argues the opposite, suggesting that their star formation driven IR and/or radio continuum emitting properties are sufficiently different to account for the large discrepancy arising from extra cooling of CR electrons by IC scattering off the CMB. Alternatively, this result may also simply suggest that, unlike their FIR output, the radio emission from both populations of SMGs is in fact dominated by AGN, or that their magnetic fields are significantly strong (i.e. $\\ga 300~\\mu$G) such that $U_{B} \\gg U_{\\rm CMB}$ at these redshifts. The latter explanation is tough to reconcile since, in the presence of such strong fields, CR electrons emitting at GHz frequencies will be cooled primarily by ionization and bremsstrahlung losses rather than by synchrotron radiation (see $\\S$2.2), unless some other physics is at work. At present it is difficult to make any strong conclusions from this comparison since it suffers from low number statistics; very few SMGs that have been identified at $z\\ga4$.   \\subsubsection{Radio Quiet QSOs at High-z} The samples of \\citet{cc04} and \\citet{ab06} consist entirely of radio-quiet quasars QSOs, a total of 6. The highest redshift QSO, J114816+525150 (hereafter, J1148+5251), is identified in Figure \\ref{fig-6} and appears to have an IR/radio ratio which is amazingly consistent with the local value. In fact, 4/6 (i.e. $\\sim$66\\%) of the QSOs have IR/radio ratios that are within 1$\\sigma$ of the canonical ratio. While these observations are suggestive that the energetics of these objects is dominated by star formation, such a conclusion is not so clear.  These galaxies may lie near the nominal IR/radio ratio due to excess emission from their AGN. Even if all of the synchrotron emission associated with ongoing star formation in these galaxies is suppressed due to other CR electron cooling processes competing successfully with respect to synchrotron radiation, the excess emission from an AGN may be enough to fortuitously put a galaxy back on the local correlation. This scenario is likely the simplest way to explain how QSO's at these redshifts have IR/radio ratios consistent with normal star-forming galaxies in the local Universe.  Alternatively, these objects would require magnetic field strengths nearing $\\ga$500~$\\mu$G levels to remain on the local FIR-radio correlation if powered by star formation alone. Taking J1148+5251 as an example; the galaxy is being observed $\\la$1~Gyr after the Big Bang. While it has been suggested that strong ($\\sim$20$~\\mu$G) turbulent fields may be present by $z\\sim10$, it is unclear how these seed fields could be amplified so rapidly and form galaxy-wide large-scale regular fields of $\\ga200-500~\\mu$G. If global star formation initiates a large-scale dynamo which acts to regulate the magnetic fields of galaxies \\citep[e.g.][]{my73,mr87,ruz88}, the timescale for this must occur rapidly as star formation is only $\\sim$400~Myr old assuming that star formation began around $z\\sim10$. The requirement for such a rapid amplification phase suggests that the large-scale field may begin to form before the initiation of star formation, during the in-fall of gas and dust \\citep[e.g.][]{ps89,rb94,hb93,kuls97}. However, these strong magnetic fields would then have to decay considerably over cosmic time to recover the typical strengths observed for galaxies in the local Universe.  % Detections of less exotic systems at these redshifts in the FIR and radio (e.g. using ALMA and the SKA) will be necessary to reconcile these outstanding questions. In fact, a comparison of the IR/radio ratios for these high-$z$ galaxies will likely be one of the only methods to test for the existence and strength of magnetic fields in the early Universe while trying to assess what role they may play in star formation.  \\subsubsection{Lyman Break Galaxies at $z\\sim3$} \\citet{cc08} recently analyzed the radio properties for a large sample of Lyman break galaxies (LBGs) at $z\\sim3$. Using a stacking analysis, and assuming the local FIR-radio correlation, they find that the radio (1.4~GHz) derived star formation rate is a factor of $1.8\\pm0.4$ times larger than that from observed UV measurements. Assuming that this difference arises from attenuation by dust at UV wavelengths, this value is significantly less than the factor of $\\sim$5 normally assumed for LBGs. They suggest that this discrepancy may be the result of depressed radio emission resulting from increased IC cooling of CR electrons off of the CMB, consistent with the analysis presented here.  The stacked 1.4~GHz flux density reported by \\citet{cc08} is $0.90\\pm0.25~\\mu$Jy, which corresponds to a rest-frame 1.4~GHz flux density of 0.68~$\\mu$Jy assuming a radio spectral index of $0.8$ \\citep{jc92}. Taking the typical size of an LBG to be $\\la$1$\\arcsec$ \\citep[e.g.][]{mg02b} results in a minimum energy magnetic field of $\\ga 8~\\mu$G \\citep{bk05}. The minimum energy magnetic field is relatively insensitive to the radio surface brightness (i.e. $B_{\\rm min} \\propto I_{\\nu}^{1/(3+\\alpha)}$); even if the reported 1.4~GHz flux density were under-predicted by a factor of 10, the corresponding magnetic field strength would only increase by a factor of $\\approx$1.8. The increased energy loss to CR electrons from IC scattering off of the CMB at $z\\sim 3$, assuming equipartition between the galaxy's magnetic and radiation field energy densities, will decrease the observed 1.4~GHz flux density by a factor of $\\sim$3.5 which is easily able to account for the discrepancy between dust attenuation factor suggested by the radio (1.4~GHz) derived star formation rate of \\citet{cc08} and the factor of $\\sim$5 normally assumed for LBGs. If the intrinsic magnetic fields of LBGs are as large as the minimum energy fields for SMGs (i.e. $\\sim 35~\\mu$G,) then the observed non-thermal radio continuum emission should still be suppressed by a factor of $\\sim$2.  \\subsection{Implications for Radio-Based Studies of Star Formation at High Redshifts \\label{sec-sfr}} While I have looked to see how the GHz emission from galaxies will be affected by increased CR electron cooling as the result of IC scattering off of the CMB at high redshifts, I  have yet to look at what this may do for higher frequency radio observations. I now consider the potential usefulness of low ($\\la$2~GHz) and high ($\\ga$10~GHz) frequency radio observations for measuring star formation over cosmic time.  \\subsubsection{At Low Frequencies: $\\la$2~GHz} As discussed above, studies of star formation at low frequencies using the SKA will be complicated by the expected evolution in the FIR-radio correlation with increasing redshift. When coupled with the additional emission from AGN, which may appear to place galaxies back on the FIR-radio correlation, it appears that using GHz observations to estimate the cosmic star formation history may lead to increasingly unreliable estimates of star formation rates with increasing redshift. Galaxies which host compact starbursts may have large magnetic fields and gas densities in which case electrons contributing to GHz emission will lose their energy primarily to ionization and bremsstrahlung processes rather than synchrotron radiation. These galaxies may move off the correlation by some fixed amount depending on how the various energy-loss terms compete, resulting in incorrect radio-based estimates of star formation rates using locally-derived calibrations. In conclusion, it seems that interpreting the non-thermal GHz measurements of high-$z$ galaxies properly may prove difficult.  \\begin{figure} \\plotone{f7.eps} \\caption{ The same as Figure \\ref{fig-3}, except that the expected emission at 10~GHz is considered. The expected sensitivity of the SKA at 10~GHz is also assumed identical to the 1.4~GHz sensitivity. Since the non-thermal contribution is almost completely suppressed at high redshift,  the 10 and 1.4~GHz emission should be detectable from the same population of galaxies at these epochs. The expected flux density range for galaxies included in the $z\\sim$ 4, 5, 6, 7, and 9 UV luminosity function studies of \\citet{rb07,rb08} is again given by the shaded region. \\label{fig-7}} \\end{figure}  \\begin{figure} \\plotone{f8.eps} \\caption{ The expected thermal fraction for observed-frame 1.4 (blue) and 10~GHz (red) observations of normal star-forming galaxies vs. redshift. Intrinsic magnetic fields of 10 ({\\it solid}-lines), 50 ({\\it dashed}-lines), and 100~$\\mu$G ({\\it dotted}-lines) are shown. By observing at 10~GHz, the observed radio emission is dominated by thermal processes, making it an ideal measure for the massive star formation rate of galaxies, beyond a redshift of $z\\sim2$ even for magnetic field strengths of $\\sim$100~$\\mu$G. \\label{fig-8}} \\end{figure}    \\subsubsection{At High Frequencies: $\\ga$10~GHz} While deep continuum studies aimed at measuring the star formation properties of high-$z$ galaxies at low frequencies (i.e. $\\la$2~GHz) are not straightforward due to large variations in the amount of non-thermal radio emission, such studies at slightly higher frequencies (i.e. $\\ga$10~GHz) are likely to be much more promising. In fact, it can be argued that deep radio continuum studies at these frequencies will provide the most accurate estimates for the star formation rate of high-$z$ galaxies. So long as average galaxy magnetic fields do not exceed mG strengths, which seems unlikely given that at such high redshifts it is unclear how such strong fields could be seeded, the observed radio emission at $\\ga$10~GHz should be dominated by thermal (free-free) radiation directly associated with Lyman continuum photons emitted by newly formed massive stars.  In Figure \\ref{fig-7} the expected observed-frame 10~GHz continuum flux densities of galaxies are plotted for a range of IR luminosities under the same assumptions that went into creating Figure \\ref{fig-3} (see Table \\ref{tbl-2}). Overplotted is the expected range in observed-frame 10~GHz flux density for the sources studied by \\citet{rb07,rb08}. It is shown that even at 10~GHz the same population of LIRGs, if they exist, should be detectable at all redshifts. Furthermore, it appears that the internal magnetic strength of a galaxy does not play as signifiant a role in a source's detectability since the thermal fraction at these frequencies is so large, and becomes increasingly so with redshift. This is illustrated in Figure \\ref{fig-8} where the expected thermal fraction for observed 1.4 and 10~GHz flux densities is plotted as a function of redshift for galaxies having 10, 50 and 100~$\\mu$G fields.  Observations at 10~GHz are dominated by thermal emission for $B\\la100~\\mu$G beyond a redshift of $\\sim$2. The same cannot be said for 1.4~GHz observations until the galaxy population being observed lies at a redshift of $z\\ga$5. At 10~GHz, the radio emission from $z\\ga5$ galaxies is likely to be $\\ga$80\\% thermal emission. Therefore, it is believed that deep continuum surveys at $\\ga$10~GHz may provide an excellent measure for the star formation activity in the highest redshift galaxies.  This conclusion will only be valid if there is a negligible amount of microwave emission associated with dust. While thermal emission from interstellar dust grains is negligible between $10-100$~GHz, so called 'anomalous' dust emission, possibly arising from a population of ultrasmall, rapidly rotating grains with a non-zero electric dipole moment, may contribute significantly at microwave wavelengths \\citep[e.g. ][]{dl98a,dl98b,ahd09}. If significant, high-frequency radio observations would likely overestimate true star-formation rates. This occurrence would also require that the observed high-$z$ galaxies have a significant amount of ultrasmall dust grains.  Comparing these observations with results from ALMA, which should be able to detect the peak FIR continuum emission from galaxies at similar high redshifts \\citep[i.e. $z\\ga$5;][]{am01}, should help disentangle this issue and result in powerful star formation and AGN diagnostics. Such tools will be critical for aiding in the interpretation of deep SKA imaging surveys which, unlike ALMA, will be able to survey large areas of the sky relatively quickly. And, as previously stated, even if the dust content in typical galaxies at these early epochs is negligible such that their FIR continuum emission is invisible to ALMA, radio continuum emission will likely be the best means to identify such objects and and quantify their star formation properties.  \\subsection{Other Physical Considerations} In the above presentation I have only considered variations to the observed IR/radio ratios which arise from an increase in the competition of energy-loss processes to primary CR electrons relative to synchrotron radiation. Alternatively, other physical processes may have their hand in driving systematic variations in a galaxy's IR/radio ratio, perhaps even keeping the correlation linear versus redshift fortuitously.  \\subsubsection{Implications for the Possible Size Evolution of Galaxies} It is possible that the suppression of a galaxy's non-thermal emission may not be as significant as previously suggested if the size of galaxies decrease with increasing redshift, resulting in larger values of $U_{B}$ per unit star formation. For instance, \\citet{rb04} suggest that the half-light radius of galaxies between $2 \\la z \\la 6$ follow an $r_{\\rm hl}\\sim(1+z)^{-1}$ relationship. If true, the $U_{\\rm CMB} \\propto (1+z)^{4}$ quenching of galaxy synchrotron emission is effectively reduced by a factor of $(1+z)^{2}$ for a constant magnetic field strength, and IC losses off the CMB will not be as significant for CR electrons. In this case such a strong departure from the canonical FIR-radio correlation may not be as easily seen. However, if galaxies do become more compact, the escape times for CR electrons to leave the galaxy will likely decrease resulting in a decrease to the observed synchrotron emission. And, if the ISM of such galaxies are typically clumpy, similar to local dwarf irregular galaxies, then escape times should decrease even more \\citep[e.g.][]{ejm08}.  For a random walk diffusion, the escape time is roughly proportional to $l_{\\rm esc}^{2}$. Assuming that $l_{\\rm esc} \\sim r_{\\rm hl} \\sim(1+z)^{-1}$, the escape time will decrease as $(1+z)^{-2}$ which should compensate for the increase to $U_{B}$ by that same factor. Furthermore, if these compact galaxies also host starbursts, advection out of each system due to a galactic wind may lead to even shorter escape times resulting in even larger IR/radio ratios at these early epochs. Therefore, even if galaxies are smaller at high-$z$, evolution in the FIR-radio correlation with redshift is still expected to be observed such that IR/radio should be larger with increasing $z$. Given that local starbursting galaxies follow the FIR-radio correlation, the suggestion for increased losses due to advection by galactic winds in such high-$z$ systems would imply that they are different than starbursts observed locally.  \\subsubsection{Early Chemical Enrichment and the First Dust Grains \\label{sec-dust}} The FIR-radio correlation may also fail at high redshifts for a variety of reasons arising from differences in a galaxy's dust emission properties. For instance, IR/radio ratios should depart from the nominal value if dust is not present, the size distribution of grains take on a different form,  or if the FIR emitting properties of dust grains are significantly different than in the local Universe. Currently, observations of quasars and luminous galaxies at high-$z$ show that these systems contain large masses of dust \\citep[e.g.][]{wang08}. While it is unlikely that these dusty systems are the representative population at early epochs, it is worth noting that the dust contained in these very young systems does not appear to be exotic enough to move systems off of the FIR-radio correlation. For example, the $z=6.42$ quasar (J1148+5251), which is plotted in Figure \\ref{fig-6} using data taken from \\citet{ab06}, has an estimated dust mass ranging between $2\\times10^{8}~M_{\\sun}$ \\citep{ed07} and  $7\\times10^{8}~M_{\\sun}$ \\citep{fb03}. Dust-to-gas mass ratios are also found to be $\\ga$0.004 which is similar to, or exceeding the, $\\sim$0.005 value of the Galaxy.  Using this QSO as an example, \\citet{bd09} states that SNe are likely responsible for producing both the metals which compose grains, as well as supernovae-condensed dust to provide surface areas for grain growth. However, \\citet{bd09} also shows that the dust in these objects are not predominantly composed of SN dust, known to have different emitting properties than that of typical dust grains in normal star-forming galaxies. Rather,  the bulk of the dust formed in these high-$z$ galaxies is attributed to grain growth occurring more rapidly than grain destruction in the ISM, resulting in large dust masses that are already in place in galaxies 400~Myr after the first episode of star formation. This estimate assumes that star formation within galaxies began at $z=10$. Given that the dust-to-gas ratios in these high-$z$ galaxies do not appear significantly different than the Milky Way value, and that the grain formation process (i.e. {\\it in~situ}) is similar to that of normal star-forming galaxies in the local Universe \\citep{bd09}, it seems reasonable to believe that the FIR dust emission from these galaxies does not drive large variations in their IR/radio ratios relative to the expected variations in their radio continuum emission.  \\subsubsection{Effect of CR Electron Acceleration Efficiencies \\label{sec-acc}} The efficiency in which CRs are accelerated in SNRs plays a significant role in determining the synchrotron emissivity from galaxies. Furthermore, the distributed re-acceleration of CR electrons arising from external interactions such as ram pressure may add significant flux to the observed radio continuum from galaxies \\citep[e.g.][]{ejm09c}. The efficiency with which the initial energy of the SNe explosion is converted into kinetic energy depends on the relative importance of radiation and adiabatic losses. In systems having extremely dense ISM, the adiabatic phase of the supernovae expansion may be halted as it expands into the ambient medium, thus reducing the amount of energy lost to adiabatic expansion. It is possible that this extra energy could be, at least partially, used in the acceleration of CRs yielding a larger amount of synchrotron emission per unit star formation than in normal galaxies. For the SMGs, whose IR/radio ratios are systematically below the FIR-radio correlation by $\\sim2-3$ (see $\\S$\\ref{sec-data}), this scenario does not appear to help explain their low ratios.  As an example, the modeling of \\citet{ed91,ed00} shows that the total acceleration efficiency for CRs increases from $\\sim$0.15$E_{\\rm SN}$ to $\\sim$0.25$E_{\\rm SN}$, where $E_{\\rm SN}=10^{51}$~erg is the total explosion energy of the SNe, when the external ISM density is increased from $\\sim$1~cm$^{-3}$ to $\\sim$10~cm$^{-3}$. The ISM densities of starbursts are likely much larger, as with the case of the example where $n_{\\rm ISM} \\sim 10^{4}$~cm$^{-3}$ was assumed, thus the evolution of SNRs will likely be different and may work to increase the synchrotron emissivity in such galaxies through an increased efficiency in particle acceleration. However, unless SMGs host starbursts that are sufficiently different than those in local starbursting galaxies, which do not appear to exhibit such variations in their IR/radio ratios, this scenario alone might not be sufficient to explain their low IR/radio ratios. At present, more simulation work on this problem is needed.  \\subsubsection{Secondary Electrons and Positrons \\label{sec-sec}} So far synchrotron radiation arising for primary CR electrons accelerated in SNRs has been considered. CR nuclei, which are accelerated in the same SNRs, will inelastically scatter off of the interstellar gas yielding $\\pi^{0}$ particles that decay into secondary $e^{\\pm}$'s and $\\gamma$-rays, and $\\pi^{\\pm}$ particles which will decay into secondary $e^{\\pm}$'s and neutrinos. These secondary $e^{\\pm}$'s will undergo the same energy-loss processes described in $\\S$2.1 and $\\S$2.2. At $\\sim$GeV energies the ratio of secondary/primary electrons in the Galaxy is roughly 2:1, and decreases to nearly 1:3 at $\\sim$3~GeV energies \\citep{tp08}. The large secondary fraction arises from the fact that there are many more CR nuclei than electrons. Consequently, diffusion models suggest that $\\sim$20-50\\% of the diffuse GHz synchrotron emission in the Galaxy arises from secondaries. However, secondary particles are produced more diffusively than primary electrons, thus their synchrotron morphology should not show any coincidence with their sites of origin. Given that the FIR and radio images of galaxies are remarkably similar \\citep[e.g.][]{ejm06,ejm08} argues that the secondaries likely contribute mostly near or around star-forming regions.  In the case of a starburst, having a much larger ISM density, the cross-section for collisions between CR nuclei and the interstellar gas is increased, thus increasing the number of $e^{\\pm}$'s for a fixed primary nuclei/electron ratio which may actually dominate the diffuse synchrotron emission. Determining the contribution of additional synchrotron emission arising by an increased population of secondary $e^{\\pm}$'s clearly goes beyond the scope of this paper, however the fact that this effect may play a role is noteworthy. The recently launched $Fermi$ $\\gamma$-ray observatory may help to shed some light on this topic.  \\subsubsection{Variations in the Initial Mass Function \\label{sec-imf}} While star-forming galaxies are currently thought to be responsible for completely reionizing the intergalactic medium (IGM) by $z\\sim6$ \\citep[e.g.][]{rb01,xf06}, the ionizing flux arising from star formation in Lyman break galaxies at similar redshifts \\citep{rb07,rb08} appears to fall a factor of $\\ga6$ below the minimum value required to maintain an ionized IGM for a given clumping factor and escape fraction under the assumption of a Salpeter stellar initial mass function (IMF). \\citet{rrc08} has shown that this discrepancy can be reconciled by flattening the stellar IMF to have a slope of $\\sim-1.7$ if reionization occurred at $z=9$. The idea of a top heavy IMF at these epochs does not seem completely inappropriate given that low metallicity environments will favor the production of more high-mass stars. In this case, one expects there to be both more SNe, increasing the amount of non-thermal synchrotron emission, as well as an increase to the ionizing photon rate per baryon, leading to an increase in the thermal radio continuum emission. It is worth noting that only a few SNe are required to enrich a galaxy from primordial values to a few tenths solar metallicity, thus most galaxies observed with ages larger than a dynamical timescale are not likely to be that metal poor \\citep[e.g.][]{rw09,hy09}.  Using the results of \\citet{rrc08}, the ionizing photon rate for a stellar IMF having a slope of $-1.7$ is a factor of $\\approx$1.7 times larger than a that for a typical Salpeter IMF slope of $-2.3$ over a mass range of $1-200~M_{\\sun}$. Similarly, for a given star formation rate, such a change to the slope of the stellar IMF should result in an increase to the core collapse supernova rate by a factor of $\\approx$2 \\citep[e.g.][]{pm98}. Thus, if the stellar IMF in these high redshift systems transitions to being top heavy, the results presented here should not be significantly affected as both the thermal and non-thermal radio continuum output should change by factors of $\\la$2. In the case that there is a significant amount of dust in these high-$z$ systems, and assuming that their non-thermal radio continuum emission is completely suppressed by the additional IC losses to CR electrons by the CMB, their IR/radio ratios should still remain similar for a varying IMF given that both the FIR and thermal radio emission are directly proportional to the ionizing photon rate. However, since the increase to the ionizing photon rate will translate into an increase to the thermal radio emission, careful analyses of radio source counts at these redshifts, along with comparison between rest-frame UV studies, may help to illuminate possible variations to the stellar IMF.  }"
    },
    "0910/0910.5291_arXiv.txt": {
        "abstract": "The internal-plateau X-ray emission of gamma-ray bursts (GRBs) indicates that a newly born magnetar could be the central object of some GRBs. The observed luminosity and duration of the plateaus suggest that, for such a magnetar, a rapid spin with a sub- or millisecond period is sometimes able to last thousands of seconds. In this case, the conventional neutron star (NS) model for the magnetar may be challenged, since the rapid spin of nascent NSs would be remarkably decelerated within hundreds of seconds due to {\\r} instability. In contrast, the {\\r}s can be effectively suppressed in nascent strange stars (SSs). In other words, to a certain extent, only SSs can keep nearly-constant extremely-rapid spin for a long period of time during the early ages of the stars. We thus propose that the sample of the GRB rapidly-spinning magnetars can be used to test the SS hypothesis based on the distinct spin limits of NSs and SSs. ",
        "introduction": "In the standard fireball model for gamma-ray bursts (GRBs; see reviews by Piran 2005; {\\mes} 2005), the GRB prompt emission is considered to arise from an internal dissipation in a relativistic expanding fireball, while the afterglow emission is produced by the deceleration of the fireball by circumburst medium, i.e., a blast wave. As understood above, it seems unlikely to directly find information about the GRB central engine, which drives the fireball, from the observed emission. However, based on some observational constraints and theoretical simulations, it is widely accepted that black holes or rapidly spinning magnetars (highly-magnetized pulsars) formed during the death of the progenitors play an important role in driving the fireball.   Furthermore, the {\\it Swift}-discovered delayed intermittent bright X-ray flares demonstrate that the GRB central objects should be still very active after the bursts (Burrows et al. 2005; Yu \\& Dai 2009). Analogically, the shallow-decaying segment remarkably appearing in the canonical X-ray light curve (Nousek et al 2006; Zhang et al. 2006) also strongly implies an energy injection into the GRB blast waves, which is quite likely to result from a long-lasting energy release from the central objects. Under these requirements, spinning-down magnetars are usually suggested as GRB central objects in the literature (e.g., Bucciantini et al. 2009; Corsi \\& {\\mes} 2009; Dai 2004; Dai et al. 2006; De Pasquale et al. 2007; Zhang \\& Dai 2008, 2009). The shallow-decaying afterglows can somewhat be regarded as an observational signature of rapidly spinning magntetars (Dai \\& Lu 1998a, b; Yu \\& Dai 2007; Zhang \\& M\\'esz\\'aros 2001).  In particular, a nearly-constant X-ray emission followed by a very steep decline with a temporal index of $\\sim9$ was reported in GRB 070110 by Troja et al. (2007). The abrupt cutoff of the X-ray light curve robustly indicates that this plateau emission is of ``internal\" origin, ruling out the external shock model. Moreover, the observed luminosity and duration of the plateau are well consistent with the parameters of a newly born magnetar. So, following Zhang \\& {\\mes} (2001), Troja et al. (2007) further argued that this puzzling plateau could be directly produced by the internal dissipation of the magnetar-driven wind at small radii, before the energy of the wind is deposited into the GRB blast wave. Then, the observed luminosity of the plateau tracking the spin-down luminosity of the GRB magnetar makes it possible to ``directly\" explore the properties of the central objects by using observed emission.  As a pioneering attempt, Lyons et al. (2009) investigated the magnetic and rotational properties of the central magnetars for ten GRB 070110-like GRBs. While the inferred magnetic field strengthes are basically typical for magnetars, the spin periods are usually found to be as short as the Kepler period, sometimes even at thousands of seconds after the GRB trigger. However, for a magnetar composing of normal neutron matter (i.e., neutron star; NS), it seems difficult to keep nearly-Kepler spin for more than one thousand seconds, because the spin of the nascent NS could be sorely restricted by some stellar instabilities, especially, {\\r} instability on focus in this {\\it letter} (S\\'a 2004; S\\'a \\& Tom\\'e 2005, 2006; Yu et al. 2009). In contrast, for a magnetar composing of strange quark matter (i.e., strange star; SS)\\footnote{According to the different physics of dense matter, a variety of models for compact stars have been constructed in the past forty years (see reviews by Weber 2005; Weber et al. 2006). Most of these models can be regarded as advanced versions of the conventional NS model. However, as an extreme case, SSs composing of strange quark matter (usually with a thin nuclear crust) are predicted to be completely different from ordinary NSs (e.g., Alcock et al. 1986).}, the \\r s-induced limit on the stellar spin is absent during the early ages of the star, because of the effective suppression of the {\\r}s by the quark matter's bulk viscosity (e.g., Wang \\& Lu 1984; Madsen 1998; Zheng et al. 2006). Therefore, we argue that {\\it a newly born magnetar, which can maintain constant rapid spin for a long period of time, could be a SS candidate}. Based on this consideration, we propose that {\\it the sample of the GRB rapidly-spinning magnetars can be used to test the SS hypothesis.}  In the next section, we briefly review the magnetar model for the internal-plateau emission, and then some magnetar parameters are derived from a small GRB sample. In Sect. 3, we analysis the {\\r} limit on the spin of NSs in the framework of a second-order {\\r} model developed by S\\'a (2004). By confronting the {\\r} limit with the GRB magnetar sample, we makes an attempt to find SS candidates. Finally, a summary and discussion are given in Sect. 4. ",
        "conclusions": "The internal-plateau emission is considered to indicate that a central rapidly-spinning magnetar is formed during some GRBs. Such a newly born magnetar is sometimes required to be able to keep a rapid spin with a sub- or millisecond period for thousands of seconds. Then, according to the spin limit of nascent NSs due to {\\r} instability, we propose a method to identify SS candidates from the GRB magnetars, since the {\\r} instability cannot operate in nascent SSs. Despite the smallness of the present GRB magnetar sample, the effectiveness of the method is somewhat exhibited by the appearance of GRB 060607A.  A few decades ago, the concept of SS was proposed formerly as a new and novel class of compact stars (Alcock et al. 1986; Haensel et al. 1986) and even as GRB central objects (Cheng \\& Dai 1996, 1998; Dai \\& Lu 1998a), based on a conjecture that the true ground state of matter is strange quark matter rather than Fe$^{56}$ (Witten 1984). In both physics and astronomy, it is of significant importance to identify or rule out SSs from observational astrophysical objects. Seasonably, a plausible method was designed by Madsen (1998) according to the distinct rotational properties of NSs and SSs against the {\\r} instability. However, for most detected Galactic millisecond pulsars, the stellar temperature is generally much lower than $\\sim10^{10}$K, in which case some complications arise from the shear viscosity due to the stellar crust (e.g., Bildsten \\& Ushomirsky 2000) and the superfluidity appearing in the stellar core. This makes it difficult to distinguish SSs from NSs within the Galactic pulsar sample (Madsen 2000). In contrast, the situation in newly born pulsars would be much clearer due to their sufficiently high temperature. Moreover, for GRB magnetars, a relatively richer sample could be selected from the huge GRB storage, in contrast to the rare observation of newly born pulsars in the Galaxy. In view of these advantages, it is worthwhile to use the newly-born GRB magnetars to test the SS hypothesis.  However, it should be notice that the uncertainties and approximations of the model still prohibit the present small GRB magnetar sample from becoming a gold sample for testing SS hypothesis. Therefore, it is demanded to make great efforts in following aspects: (1) a thorough investigation on magnetar wind. Here we would like to point out that the angle of the magnetar wind is probably different from that of the GRB ejecta, in respect that the GRB ejecta could be driven by an hyperaccretion onto the magnetar rather than the magnetar wind (Zhang \\& Dai 2008, 2009). (2) finding other constraints on the wind angle. For example, in order to build up sufficient dynamo action responsible for the intense magnetic fields, the initial spin periods of magentars could be required to be $\\leq10$ ms (Lyons et al. 2009; Usov 1992), which leads to $\\theta_{w}>16^{\\circ}$; (3) accumulating a much larger number of GRB magnetars."
    },
    "0910/0910.5258_arXiv.txt": {
        "abstract": "{Various authors have suggested that the gamma-ray burst (GRB) central engine is a rapidly rotating,  strongly magnetized,  $(\\sim 10^{15}-10^{16}$ G) compact object. The strong magnetic field can accelerate and collimate the relativistic flow and the rotation of the compact object can be the energy source of the GRB. The major problem in this scenario is the difficulty of finding an astrophysical mechanism for obtaining such intense fields. Whereas, in principle, a neutron star could maintain such strong fields, it is difficult to justify a scenario for their creation. If the compact object is a black hole,  the problem is more difficult since,  according to general relativity it has ``no hair\" (i.e., no magnetic field).   Schuster, Blackett, Pauli,  and others have suggested that a rotating neutral body can create a magnetic field by non-minimal gravitational-electromagnetic coupling (NMGEC). The Schuster-Blackett form of NMGEC was obtained from the Mikhail and Wanas's  tetrad theory of gravitation (MW). We call the general theory NMGEC-MW. We investigate here the possible origin of the intense magnetic fields $\\sim 10^{15}-10^{16}$  G in GRBs by NMGEC-MW. Whereas these  fields are difficult to explain astrophysically, we find that they are easily explained by NMGEC-MW. It not only explains the origin of the $\\sim 10^{15}-10^{16}$G fields when the compact object is a neutron star, but also when it is a black hole.} ",
        "introduction": "Gamma-ray bursts (GRBs) are short and intense pulses of soft $\\gamma$ rays. The bursts last from a fraction of a second to several hundred seconds. Magnetic fields in GRBs  are believed to play  a  crucial role in the internal engine as a possible way to power and collimate  the relativistic outflow   \\cite{pir05b,lee00}. Energy considerations  require  extremely large magnetic fields on the order of $10^{15}$G  \\cite{pir05,uso92,lee00}. The relativistic outflow is a Poynting flux (with negligible baryon content) \\cite{pir05}.  Usov \\cite{uso92} suggested that GRBs arise during the formation of rapidly rotating  highly magnetized neutron stars.  Very strong magnetic fields could  form and the pulsar could produce a relativistic Poynting flux flow.  Various authors have investigated this magnetar-GRB connection \\cite{tho94, kat97, klu98, whe00, tho04, met07, buc06, kom07, buc07, yu07}.   Thompson \\cite{tho07} reviewed the arguments for proto-magnetars as possible GRB engines. Long duration GRBs ($\\sim$ several hundred seconds) are, in general, attributed to core-collapse supernova. A core-collapse supernova leaves behind a hot proto-neutron star that cools on the Kelvin-Helmholtz time scale $\\tau_{KH}~ (\\sim 10-100 s$) by radiating neutrinos. Assuming that neutrino absorption on free nucleons dominates heating, a characteristic thermal pressure is $\\sim 3 \\times 10^{28} ergs ~cm^{-3}$. Setting $B^{2}/8\\pi$ equal to this pressure requires $B\\sim 10^{15}$ G. Thus if $B$ is involved in collimating and accelerating the relativistic jet, a field $B\\gtrsim 10^{15}$ G is necessary. Like beads on a wire, the magnetic field lines force the plasma around the proto-magnetar into co-rotation with the stellar surface out to the Alfven surface, where the magnetic energy density equals the kinetic energy density of the outflow. A neutron star with a millisecond spin period has a reservoir   of rotational energy of $\\sim 2 \\times 10^{52} P^{-2} ergs$, where $P$ is the period of rotation of the neutron star in milliseconds. Thus, the rotational energy of the proto-magnetar   can be tapped and could produce a GRB if its rotational period is $\\sim$ milliseconds.  In summary, the  reasons for considering millisecond proto-magnetars as the central engine of long-duration GRB are: \\begin{enumerate} \\item the reservoir of rotational energy is in the range required to power GRBs; \\item the characteristic time of the long-duration GRBs is on the order  of the Kelvin-Helmholtz cooling time, $\\tau_{KH}$; \\item there is an observational connection between core collapse supernovae and long duration GRBs; and \\item the magnetic field  of a magnetar $\\gtrsim 10^{15}$ G is strong enough to accelerate and collimate the outflow. \\end{enumerate}  The  major problem of  the magnetar model is the origin of the $\\sim 10^{15}-10^{16}$ G fields. As noted by Spruit \\cite{spr08}, inheritance of the magnetic field from a main sequence progenitor and dynamo action at some stage of the progenitor is not sufficient to explain the fields of magnetars. A runaway phase of exponential growth would be needed to achieve sufficient field amplification. The  magnetorotational instability is a possibility, but the magnetic  field would probably decay rapidly by magnetic instabilities.    In a core collapse supernova, a black hole can be formed instead of a neutron star. The justification of the presence of strong magnetic fields in that  case is more difficult since, according to general relativity,  black holes have no magnetic fields.  The Schuster-Blackett (S-B) conjecture \\cite{sch11, bla47}  suggests a gravitational origin of the magnetic fields in rotating neutral bodies, i.e., a  non-minimal gravitational-electromagnetic coupling (NMGEC).  An early attempt to make a theory that encompasses  the S-B conjecture  was made  by  Pauli  \\cite{pau33}. Latter, attempts were made by   Bennet \\textit{et al.} \\cite{ben49}, Papapetrou \\cite{pap50} , Luchack  \\cite{luc52},  and Barut \\cite{bar85}. The majority of these studies were based on the five-dimensional Kaluza-Klein formalism. This formalism was used in order  to describe a unified theory of gravitation and electromagnetism  in such a way so  that the S-B conjecture is obtained. Opher and Wichoski \\cite{oph97} applied the S-B conjecture for the   origin of magnetic fields in rotating galaxies and proposed that the $B \\sim 10^{-6}-10^{-5}$ G magnetic fields in spiral galaxies are directly obtained from NMGEC. Mikhail \\textit{et al.} \\cite{ mik95} showed that Mikhail and Wanas tetrad theory of gravitation (MW) \\cite{mik77} predicts the S-B conjecture of NMGEC. We call this the NMGEC-MW theory.  We investigate here  the possibility that NMGEC-MW is the origin of the intense magnetic fields $10^{15}-10^{16}$ G, connected with  GRBs produced by rotating neutron stars or black holes. In section 2,  we discuss NMGEC and in section 3,  NMGEC-MW. In section 4,  we apply NMGEC-MW to GRBs. Our conclusions and discussion are presented in section 5. ",
        "conclusions": "Magnetic fields $\\simeq 10^{15}-10^{16}$ G have been suggested to be connected with central engines of GRBs. We considered the possibility that GRBs are powered by a rapidly rotating highly magnetized central compact object, a black hole or a neutron star.   Since these magnetic fields are  difficult to be obtained   by  astrophysical mechanisms, we considered the possibility that the fields are created by NMGEC-MW.  For  GRBs a magnetic field $\\sim 10^{15}-10^{16}$ G is required to produce the Poynting flux needed to supply the energy observed within the  required time.  The  fields  predicted by  NMGEC-MW in Eq.  (\\ref{Bgama}) for a black hole and Eq. (\\ref{bpfim}) for a neutron star,  are   in agreement with models  requiring central engines with $\\gtrsim 10^{15}$ G fields. We conclude that NMGEC-MW is a possible mechanism for  creating $\\gtrsim 10^{15}$ G fields in the central engine of GRBs due to a rotating neutron star or black hole.  From the results of our paper, one might believe that stellar mass black holes lose the rotation generated in their formation very quickly due to magnetic dipole emission from the strong magnetic field generated by NMGEC-MW. We might then conclude that only the youngest stellar black holes observed have non-zero rotation. However, stellar mass black holes are only observed in accreting binaries. Accretion from the stellar companion makes the black hole visible in X-rays. The accretion is able to spin-up the black holes. Thus, old black holes, as well as young black holes,  can have non-zero rotation.    \\clearpage"
    },
    "0910/0910.0966_arXiv.txt": {
        "abstract": "PKS 1510$-$089 is a powerful Flat Spectrum Radio Quasar at z=0.361 with radiative output dominated by the $\\gamma$-ray emission. In the last two years PKS 1510$-$089 showed high variability over all the electromagnetic spectrum and in particular very high $\\gamma$-ray activity was detected by the Gamma-Ray Imaging Detector on board the AGILE satellite with flaring episodes in August 2007 and March 2008. An extraordinary activity was detected in March 2009 with several flaring episodes and a flux that reached 500 $\\times$ 10$^{-8}$ photons cm$^{-2}$ s$^{-1}$. Multiwavelength observations of PKS 1510$-$089 seem to indicate the presence of Seyfert-like features such as little blue bump and big blue bump. Moreover, X-ray observations suggested the presence of a soft X-ray excess that could be a feature of the bulk Comptonization mechanism. We present the results of the analysis of the multiwavelength data collected by GASP-WEBT, REM, {\\it Swift} and AGILE during these $\\gamma$-ray flares and the theoretical implications for the emission mechanisms. \\vspace{-0.5cm} ",
        "introduction": "Blazars are the most extreme class of Active Galactic Nuclei (AGNs) characterized by the emission of strong non-thermal radiation across the entire electromagnetic spectrum, from radio to very high $\\gamma$-ray energies. The typical observational properties of blazars include irregular, rapid and often very large variability, apparent super-luminal motion of the jet at VLBI scales, flat radio spectrum, high and variable polarization at radio and optical frequencies. These features are interpreted as the result of the emission of electromagnetic radiation from a relativistic jet that is viewed closely aligned to the line of sight, thus causing strong relativistic amplification (Blandford $\\&$ Rees 1978; Urry $\\&$ Padovani 1995).  The EGRET instrument onboard $Compton$ $Gamma$-$Ray$ $Observatory$ detected for the first time strong and variable $\\gamma$-ray emission from blazars in the MeV-GeV region. In conjunction with ground-based Cherenkov telescopes and coordinated multiwavelength observations, this provided evidence that the Spectral Energy Distributions (SEDs) of blazars are typically double-humped with the first peak occurring in the IR/optical band for the so-called $red$ $blazars$ (including Flat Spectrum Radio Quasars, FSRQs, and Low-energy peaked BL Lacs, LBLs) and at UV/X-rays for the so-called $blue$ $blazars$ (including High-energy peaked BL Lacs, HBLs). The first peak is commonly interpreted as synchrotron radiation from high-energy electrons in a relativistic jet. The SED second component, peaking at MeV-GeV energies in $red$ $blazars$ and at TeV energies in $blue$ $blazars$, is commonly interpreted as inverse Compton (IC) scattering of seed photons, internal or external to the jet, by relativistic electrons (Ulrich et al., 1997), although other models have been proposed (see e.g., B\\\"ottcher 2007 for a recent review). Blazars with greater bolometric luminosity have smaller peak frequencies and ``redder'' SEDs, while blazars of lower bolometric luminosity have higher peak frequencies and then are ``bluer'' (Fossati et al., 1998). This spectral sequence was interpreted by Ghisellini et al.~(1998)  as consequence of the different radiative cooling suffered by the emitting electrons of blazars of different power.  With the detection of several blazars in the $\\gamma$-rays by EGRET (Hartman et al., 1999) the study of this class of objects has made significant progress. In fact, considering that the large fraction of the total power of blazars is emitted in the $\\gamma$-rays, information in this band is crucial to study the different radiation models. The interest in blazars is now even more renewed thanks to the simultaneous presence of two $\\gamma$-ray satellites, AGILE and $Fermi$, and the possibility to obtain $\\gamma$-ray observations simultaneously with multiwavelength data collected from radio to TeV energies will allow us to reach a deeper insight on the jet structure and the emission mechanisms at work in blazars. ",
        "conclusions": ""
    },
    "0910/0910.2146_arXiv.txt": {
        "abstract": "{} {We report new VLBA polarimetric observations of the compact steep-spectrum (CSS) quasar 3C\\,147 (B$0538+498$) at 5 and 8.4~GHz. } {By using multifrequency VLBA observations, we derived milliarcsecond-resolution images of the total intensity, polarisation, and rotation measure distributions, by combining our new observations with archival data.} {The source shows a one-sided structure, with a compact region, and a component extending about 200 mas to the south-west. The compact region is resolved into two main components with polarised emission, a complex rotation measure distribution, and a magnetic field dominated by components perpendicular to the source axis.} {By considering all the available data, we examine the possible location of the core component, and discuss two possible interpretations of the observed structure of this source: core-jet and lobe-hot spot. Further observations to unambiguously determine the location of the core would help distinguish between the two possibilities discussed here.} ",
        "introduction": "\\label{sec:intro}  Compact steep-spectrum (CSS) radio sources are scaled-down versions of large-sized double sources with linear sizes $\\leq 20\\,h^{-1}$\\,kpc \\footnote {$q_0=0.5$ and $H_0=100\\,h\\,{\\rm km}\\,{\\rm s}^{-1}\\,{\\rm Mpc}^{-1}$} and steep high-frequency radio spectra ($\\alpha >0.5$ \\footnote{we assume ${\\rm S}_{\\nu}=\\nu^{-\\alpha}$}). It is now generally believed that most CSS sources are {\\it young}, of ages $<10^{3-5}$\\,yr, whose radio lobes have had insufficient time to grow to kilo-parsec scales ~\\citep{Fanti95}. Subgalactic in size, they reside deep inside their host galaxies and are largely confined to the Narrow Line Region (NLR) with its relatively large column density of ionised plasma. This magnetised thermal plasma has different indices of refraction for the two circular polarised modes of the synchrotron emission of these sources. This rotates the orientation of the linearly polarised components of the transmitted radiation. The amount of the rotation is given by $\\Delta\\chi=812\\times\\lambda^2\\,\\int{n_{\\rm e} B_{||} d\\ell}$, where $\\lambda$ is the wavelength in cm, $n_{\\rm e}$ is the electron density  of the medium in cm$^{-3}$, $B_{||}$ is the component of the magnetic field along the line of sight in $\\mu$G,  and $\\ell$ is the geometrical depth of the medium along the line of sight in kpc.  Since the rotation is proportional to $\\lambda^2$, it can be determined from observations at a number of wavelengths. The Faraday {\\it Rotation Measure}, {\\it RM}, is the amount of rotation expressed in rad$\\,$m$^{-2}$. Even for moderate magnetic field intensity for the foreground screen, significant Faraday rotation effects external to the radio sources are to be expected, which can completely depolarise sources at long wavelengths. The comparison of polarisation images over a range of wavelengths, preferably with similar resolutions, is then an important diagnostic of the physical conditions within and around CSS radio sources. CSS radio galaxies show very low, or even no, linear polarisation at centimetre wavelengths. CSS quasars are instead found to have polarisation percentages as high as 10\\% above 1~GHz \\citep{Saikia85, Saikia87, Fanti04, Rossetti08}.  A minority of CSSs exhibit complex or highly-asymmetric structures. Observations indicate that intrinsic distortions are caused by interactions with a dense, inhomogeneous gaseous environment ~\\citep{Mantovani94}. This view is supported by the most distorted and complex structures being found in jet-dominated objects with very weak cores ~\\citep{Mantovani02}. Furthermore, an increased asymmetry in terms of intensity, arm ratio, spectral index, and polarisation with decreasing source linear size suggests that they are expanding through the dense inhomogeneous interstellar medium (ISM) of their host galaxies ~\\citep{Sanghera95, Saikia01, Saikia03, Rossetti06}.  Sub-arcsec polarimetry provided evidence of the interaction of components of CSS sources with dense clouds of gas ~\\citep{Nan99}. According to ~\\citet{Junor99a}, the CSS source \\object{3C\\,147} has the expected signatures of a jet colliding with a cloud of gas on the southwestern side of the galaxy. Optical observations also show evidence of such interactions. \\citet{GW94} found that CSS sources have relatively strong, high equivalent width, high excitation line emission, with broad, structured [OIII]\\,$\\lambda\\,5007$ profiles, which they interpret as evidence of strong interaction between the jet and the ISM.  The number of CSSs for which detailed radio polarisation information is available remains small. To improve the statistics, we are conducting a series of observations to image examples with moderate degrees of polarised emission and clear signatures of interaction with their environment, namely fractional polarisations that decrease with increasing wavelength, and $RM>450$\\,rad\\,m$^{-2}$ in the source rest frame. Results for the first two CSS quasars observed in this ongoing program, \\object{B$0548+165$} and \\object{B$1524-136$}, are available in  ~\\citet{Mantovani02}.  We present results for \\object{3C\\,147} (\\object{B$0538+498$}), imaged with milliarcsecond resolution via full-Stokes VLBA observations. This CSS radio source is associated with a QSO of $m_{v}=16.9$ and $z=0.545$ \\citep{Spinrad85}. It has a total angular extent of $0\\farcs 6$, corresponding to a projected linear size of $\\approx 2.2\\,h^{-1}\\,{\\rm kpc}$. \\object{3C\\,147} shows a very large integrated {\\it RM} of $\\approx -1500\\,{\\rm rad}\\, {\\rm m}^{-2}$ in the observer's frame \\citep{Inoue95}, implying that the radio source is surrounded by a dense ionised medium. On the sub-arcsecond scale \\object{3C\\,147} exhibits a very asymmetric double-lobed radio structure with a strong, steep-spectrum, central component and asymmetric components of even steeper spectrum situated to the southwest and north \\citep{A-G95, Ludke98, Junor99a}. An HST optical study of the emission-line region of CSS radio sources ~\\citep{Axon00} shows that the emission-line morphology ([OIII]\\, $\\lambda\\,5007$) of \\object{3C\\,147} is dominated by an unresolved central component, embedded in a more diffuse region, which extends over $\\approx 1\\farcs2$ and is well aligned with the double radio structure observed on the sub-arcsecond scale.  \\citet{Junor99a} imaged the polarised emission of \\object{3C\\,147} using A-array VLA observations at X- and U-bands. They measured significant differential Faraday rotation between the northern and the main (central+south-western) components situated on opposite sides of the quasar, and a strong east-west gradient in the polarisation percentage of the main component, whose emission depolarises strongly towards lower frequencies.  They interpreted these results in terms of an asymmetric environment with large amounts of thermal gas on the western side.  Very Long Baseline Interferometry (VLBI) observations have been analysed to indicate that the central component contains a compact core with a jet emerging from it toward the southwest \\citep{Alef90, Nan00}. Part of this jet yields evidence of mild superluminal motion: \\citet{Alef90} reported $\\approx 1.3\\,c$, and \\citet{Nan00} inferred $\\approx 1.5\\,c$.  In this paper, we present multi-frequency VLBA (plus a single VLA antenna) polarisation observations at 5 and 8.4~GHz. In Sect. \\ref{sec:observation} we summarise the observations and data processing. Section \\ref{sec:results} describes new information obtained about the morphological and polarisation properties of \\object{3C\\,147}. Discussion and conclusions are presented in Sects. \\ref{sec:discussion}  and \\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  By analysing new polarimetric observations at C- and X-band, and re-analysing previous VLBA observations, we have been able to image the milliarcsecond structure of the quasar \\object{3C\\,147}. We have identified a complex central region that is dominated by two bright components, A and B, with a weak, compact component to the northeast, A$_0$, and a further extended component to the southwest, which is imaged out to a distance of 200\\,mas.  Polarised emission has been detected for both of the compact components, A and B, at C- and X-bands. There are indications of polarised flux-density variability in component A on a timescale of several years. The polarised emission from component B divides into two independent regions in our observations. The two compact components also show very high values of  {\\it RM} in the rest frame of the source, these being consistent with the rest-frame value reported for the entire main component by \\citet{Junor99a} of $RM =-3140\\pm630$\\,rad\\,m$^{-2}$.  The intrinsic orientation of the magnetic field in the central region of \\object{3C\\,147}  follows an arc-shape structure at the southern end of the polarised region. The magnetic-field orientation is almost perpendicular to the source major axis. Similar results were obtained by \\citet{Junor99a} at sub-arcsecond resolution. The separation between the two central components of the source seems to be increasing with an apparent velocity of $1.2\\pm0.4\\, c$.  A straightforward interpretation of all the observational data for this radio source is not yet possible. We have discussed the possible location of the core and explored two possible scenarios: a core-jet source, and a lobe-hot spot structure. Further polarisation images, and a more thorough search for the source core, should help us to discriminate between these two scenarios."
    },
    "0910/0910.4410_arXiv.txt": {
        "abstract": "We examine the fate of fast electrons (with energies $E>10 \\eV$) in a thermal gas of primordial composition.  To follow their interactions with the background gas, we construct a Monte Carlo model that includes: (1) electron-electron scattering (which transforms the electron kinetic energy into heat), (2) collisional ionization of hydrogen and helium (which produces secondary electrons that themselves scatter through the medium), and (3) collisional excitation (which produces secondary photons, whose fates we also follow approximately).  For the last process, we explicitly include all transitions to upper levels $n \\le 4$, together with a well-motivated extrapolation to higher levels.  In all cases, we use recent calculated cross-sections at $E<1 \\keV$ and the Bethe approximation to extrapolate to higher energies.  We compute the fractions of energy deposited as heat, ionization (tracking HI and the helium species separately), and excitation (tracking HI \\lya separately) under a broad range of conditions appropriate to the intergalactic medium.  The energy deposition fractions depend on both the background ionized fraction and the electron energy but are nearly independent of the background density.  We find good agreement with some, but not all, previous calculations at high energies.  Electronic tables of our results are available on request. ",
        "introduction": "\\label{intro}  Fast electrons, typically generated by high-energy photons or cosmic ray collisions, are crucial for a wide range of astrophysical problems.  For example, cosmic ray ionization is important for the thermal balance of interstellar clouds \\citep{spitzer69}, electrons produced by the Comptonization of $\\gamma$-ray photons affect the early history of supernova remnants \\citep{xu91b}, and broad-line emission regions around quasars are likely ionized by hard photons from those sources, making the fate of the liberated electrons of paramount importance \\citep{shull85}.  As such, there have been numerous studies of the interactions between these electrons and a background gas, through the processes of collisional excitation, ionization, and electron-electron scattering.  These calculations have included both analytic approaches \\citep{spitzer68, spitzer69, jura71, bergeron73, xu91} and numerical explorations \\citep{habing71, shull79, shull85, valdes08}.  Most have focused on interactions with atomic or ionic gases, relevant especially to low-density material in the interstellar (or intergalactic) media.  Others have considered the additional effects of molecules (e.g., \\citealt{dalgarno99}).  Recently, the fate of X-rays in the high-redshift intergalactic medium (IGM) has become an important question.  Before the reionization of HI, ultraviolet ionizing photons are trapped near their sources, but X-rays can travel much larger distances through the IGM.  As such, they are thought to provide the most important radiation background for the bulk of the IGM, slowly ionizing the mostly  neutral gas far from stars.  However, most of their energy is deposited indirectly, through the fast electrons generated by photo-ionization.  For example, the electrons produced by X-rays from the first hard sources (quasars, supernova remnants, or even Population III stars; \\citealt{oh01, venkatesan01}) are most likely responsible for heating the IGM to $T \\sim 1000 \\kel$ before reionization  \\citep{kuhlen05-sim, furl06-glob}.  This heating and ionization could influence future generations of structure (e.g., \\citealt{ricotti02, oh03-entropy}) and it has important observational consequences.  For example, it can dramatically change the highly-redshifted 21 cm signal from the early IGM \\citep{furl06-review, kuhlen06-21cm, pritchard07}, and it could even provide a signature of exotic physics during the Dark Ages before the first sources appear \\citep{chen04-decay, furl06-dm, mapelli06, valdes08}.  Moreover, HI \\lya photons produced as the secondary electrons collisionally excite atoms in the background gas can greatly affect the 21 cm signal by changing the hyperfine level populations of HI \\citep{chuzhoy06-first, pritchard07, chen08}, and the secondary ionizations may play an important role in the HI (and HeI) reionization of the IGM \\citep{ricotti04_a, ricotti05, volonteri09}.  Here we revisit this problem using a Monte Carlo model (described in \\S \\ref{mc}) that incorporates the most recent cross-sections for interactions between X-rays and primordial gas.  We review the relevant physics of collisional ionization, excitation, and electron scattering in \\S \\ref{ion}-\\ref{heat} and comment on some of our approximations in \\S \\ref{ignored}.  We present our principal results, together with a comparison to previous work, in \\S \\ref{results} and a simple example of their utility for heating of the high-redshift IGM in \\S \\ref{example}.  Finally, we conclude in \\S \\ref{disc}. ",
        "conclusions": "\\label{disc}  Using a Monte Carlo model, we have re-examined the fate of fast electrons scattering through a background gas of primordial origin.  We included electron-electron scattering as well as collisional ionization and excitation of HI, HeI, and HeII, explicitly tracking all levels up to $n=4$ and using an analytic extrapolation to higher levels.  We separately followed all excitations producing HI \\lya photons, which can be important in modeling the observable properties of the IGM at high redshifts (see \\citealt{kuhlen06-21cm} and \\citealt{furl06-review}, for example) and have not been explicitly tracked previously except at the highest energies.  We used recent calculations of ionization and excitation cross-sections at $E<1 \\keV$ and extrapolated to higher energies using the Bethe approximation.  In highly neutral gas ($x_i \\la 10^{-3}$), we found that $\\sim 20\\%$ of the electron energy is deposited as heat, with the remainder split roughly equally between ionization and excitation.  In this regime, the results are not strongly sensitive to $x_i$, at least at high energies, because most of the heating comes from secondary electrons, with energies below $10 \\eV$.  At higher $x_i$, the heating fraction rises rapidly, exceeding $\\sim 65\\%$ by $x_i \\sim 0.1$.  We find that the excitation and ionization energy deposition rates are always comparable, and that $\\sim 80\\%$ of the excitation energy goes into HI \\lya regardless of electron energy and $x_i$.  Although our calculations used parameters appropriate to the low-density IGM, our results may also be applied to denser systems, because the density only enters through the Coulomb logarithm affecting the electron-electron scattering rate.  Varying the density of target atoms by many orders of magnitude only affects the energy deposition fractions by a few percent in absolute terms.  We also note that, when collective plasma effects are included, the background temperature becomes irrelevant for the electron energies under consideration \\citep{schunk71,xu91}.  For the most part, our results agree with previous estimates from \\citet{shull85} and \\citet{xu91}, although we have found some discrepancies with the commonly-used fitting formulae from the former.  In general, at high energies their results slightly underestimate the importance of collisional excitation but are otherwise accurate.  At lower energies, the differences in cross-sections become more important and the discrepancies increase (at least with reference to the fitting formulae of \\citealt{ricotti02}).  However, our results show substantial differences with those of \\citet{valdes08}, who examined the high-energy limit with a Monte Carlo model similar to ours.  These discrepancies are especially large at moderate and high ionized fractions, where we find substantially more heating and less ionization and excitation.  We also find that a higher fraction of excitation energy is deposited in HI \\lya photons.  The latter is probably due to our different treatments of collisional excitation (in particular, their neglect of excitations to states other than the $np$ sublevels), but the source of the former is unclear.  In any case, we advocate interpolation of the exact results when high accuracy is necessary, especially because the energy dependence is quite significant.  Our detailed numerical results are available upon request, including tables for the energy deposition fractions at a variety of ionized fractions and C code to interpolate to arbitrary ionized fractions and electron energies.  \\vskip 0.5in  We thank M. Valdes for helpful comments. SRF was partially supported by the NSF through grant AST-0829737, by the David and Lucile Packard Foundation, and by NASA through the LUNAR program. The LUNAR consortium (http://lunar.colorado.edu), headquartered at the University of Colorado, is funded by the NASA Lunar Science Institute (via Cooperative Agreement NNA09DB30A) to investigate concepts for astrophysical observatories on the Moon.  SJS was supported by a Research Experiences for Undergraduates grant at UCLA, NSF PHY-0552500.  We thank Agner Fog for making a public version of the Mersenne Twister algorithm available and Igor Bray and Yuri Ralchenko for making the CCC database available electronically."
    },
    "0910/0910.1832_arXiv.txt": {
        "abstract": "At present, microlensing light curves from different telescopes and filters are photometrically aligned by fitting them to a common model.  We present a second method based on photometry of common field stars.  If two spectral responses are similar (or the color of the source is known) then this technique can resolve important ambiguities that frequently arise when predicting the future course of the event, and that occasionally persist even when the event is over.  Or if the spectral responses are different, it can be used to derive the color of the source when that is unknown. We present the essential elements of this technique and apply it to the case of MOA-2007-BLG-192, an important planetary event for which the system  may be a terrestrial planet orbiting a brown dwarf or very low mass star. The refined estimate of the source color that we derive here $V-I=2.36\\pm 0.03$ will aid in making the estimate of the lens mass more precise. ",
        "introduction": "\\label{sec:intro}}  The technical question most frequently asked of microlensers is ``How do you align different data sets''?  The standard answer is that data are aligned via a common microlensing model.  That is, there is a microlensing magnification model $A(t)$, and the fluxes observed at the $i$th observatory are fit to $F_i(t) = f_{s,i} A(t) + f_{b,i}$, where $f_{s,i}$ is the instrumental source flux for that observatory and $f_{b,i}$ is blended light that does not participate in the event. This approach is very powerful: it allows microlensers to work with uncalibrated data, in non-standard filters, and to use difference image analysis (DIA) photometry without worrying about the {\\it flux} zero point (which is simply absorbed into $f_b$).  These advantages are all very important because microlensing data must often be analyzed very quickly, even when they come from unexpected quarters.  Nevertheless, there are instances when an alternative alignment method would be helpful.  As we outline below (and discuss more extensively in Section~\\ref{sec:discuss}) the need for such an alternative method has been at least subconsciously apparent for several years.  However, we were motivated to systematically develop it by the problem of measuring the source color for the interesting planetary event,  MOA-2007-BLG-192.  When microlensing planet searches were first proposed \\citep{liebes64,mao91,gouldloeb92}, there was no expectation that the planet masses, distances, planet-star physical separations, or orbital motion would be determined on an individual basis.  Rather, it was thought that the quantities that could be measured were the planet/star mass {\\it ratio}, $q$, and the planet-star projected separation $d$ in units of the Einstein radius, $\\theta_\\e$. Physical information about the planets would be restricted to statistical statements made about the ensemble of detections.  In fact, of the 9 microlensing planets published to date, the masses, projected separations, and distances are measured for four \\citep{ob03235,bennett235,ob05071,ob05071b,ob06109}, and are at least partly constrained for the rest \\citep{ob05390,ob05169,mb07192,mb07400,mb08310}. The primary reason for this turnabout is that, in strong contrast to garden-variety microlensing events, planetary events usually give rise to measurable finite-source effects. These then permit determination of \\begin{equation} \\rho\\equiv{\\theta_*\\over\\theta_\\e}, \\label{eqn:rho} \\end{equation} the ratio of the angular source size to the angular Einstein radius. If $\\theta_*$ can be determined, then one can measure $\\theta_\\e$ \\begin{equation} \\theta_\\e = \\kappa M \\pi_\\rel, \\qquad \\kappa \\equiv {4 G \\over c^2\\,\\rm AU}\\sim 8.1\\,{{\\rm mas}\\over M_\\odot}, \\label{eqn:thetae} \\end{equation} where $M$ is the lens (host star) mass and $\\pi_\\rel$ is the lens-source relative parallax.  If one can then obtain a constraint on another combination of lens mass and distance, from measuring e.g., the so-called ``microlens parallax'' \\citep{gould00}, the flux from the lens \\citep{han05,bennett07}, or astrometric offsets \\citep{bennett235,ob05071b}, then one can solve for both $M$ and $\\pi_\\rel$, and so obtain the planet mass (since $q$ is usually well-measured), as well as the distance to the lens. (Since the source is almost certainly in the Galactic bulge, $\\pi_\\rel$ directly yields the lens distance.)\\ \\ The lens distance, $D_{\\rm L}$, then allows one to infer the projected separation $r_\\perp = D_{\\rm L}\\theta_{\\rm E} d$. Even if no other constraints are obtained, however, measurement of $\\theta_{\\rm E}$ still yields the product $M\\pi_\\rel = \\theta_\\e^2/\\kappa$, which then provides statistical constraints on the properties of the lens that are far better than if $\\theta_{\\rm E}$ is not measured.  Hence, there is a high premium on measuring $\\theta_\\e$ during planetary microlensing events.  The standard method for doing this is to measure the dereddened color $(V-I)_0$ and magnitude $I_0$ of the source during the event.  The dereddened color gives the surface brightness \\citep{kervella04}, and the dereddened flux then gives the angular source size \\citep{yoo04}. In fact, to a good approximation, all one really needs is the {\\it instrumental} magnitude (which is automatically returned by the light curve model) and the {\\it instrumental} color (which can be determined even without a model, just assuming that one has near-simultaneous photometry in $V$ and $I$ at several different magnification levels.  The source can then be placed on an instrumental color-magnitude diagram (CMD) and compared to the position of the red giant clump, whose dereddened color and magnitude are known fairly well.  Since the source suffers nearly the same extinction as the clump, one can directly determine $(V-I)_0$ and $I_0$ from such a diagram.  And therefore, microlensing planet hunters always try to obtain $V$-band measurement while the source is significantly magnified, to supplement the routinely-obtained $I$-band data. In fact, they try to obtain $H$-band data as well, since a 3-band $VIH$ determination can yield an even more precise measurement of $\\theta_*$ \\citep{bennett10,ob07224}.  Nevertheless, for a variety of reasons, including bad weather as well as the general chaos that is an indelible part of chasing after high-magnification microlensing events, sometimes these data are not taken or are not of adequate quality.  In the case of MOA-2007-BLG-192, no $V$-band data were taken simply because the event was not recognized as being sensitive to planets until after peak, and was not recognized as containing a planet until it had returned to baseline.  \\citet{mb07192} were nevertheless able to make a rough estimate of the source color by measuring the source magnitude (as described above) and assuming that it is a main-sequence star in the Galactic bulge.  While these assumptions are not unreasonable, they lead to fairly large errors, and could in principle fail catastrophically if the source happened, e.g., to be in the Sagittarius Dwarf galaxy. Hence it would certainly be better to have a measured color than an estimated one.  This is particularly true because the planetary system detected in this event is quite interesting.  The most favored model is for a brown-dwarf host with a few-Earth-mass planet.  Substantial work will be required to obtain the necessary constraints to confirm or reject this model \\citep{mb07192}, but a color measurement (and so a measurement of $\\theta_*$) would certainly be a step in the right direction.  Here we present a general method of obtaining such post-facto color measurements and apply it to MOA-2007-BLG-192.  We find that it is somewhat redder than originally estimated, but well within the previous (appropriately generous) error bar.  The method can potentially be applied to obtain colors of other interesting microlensed sources.  Perhaps even more important, it can be inverted to align data sets during the early phases of microlensing events when the model is poorly constrained, thus enabling much better real-time predictions, which are crucial to organizing observations.  In some infrequent but nonetheless important cases, the relative flux normalizations from different observatories remain different for different event models, even after the event is over.  Finally, it can be applied to obtain ``microlens parallaxes'' by aligning space-based and ground-based photometry.  The method we describe here can be used to untangle all these cases as well.   { ",
        "conclusions": "\\label{sec:conclude}}  We have presented a method for aligning microlensing light curves from different observatories that is independent of the standard method, which is based on fitting to a common model.  We were initially motivated to develop this technique in order to measure the source color for the archival event MOA-2007-BLG-192, which is a candidate brown-dwarf lens hosting a terrestrial planet.  We succeeded in measuring this color within $\\sigma(V-I) = 0.03$, which will aid in future efforts to characterize this planetary system.  We have argued that the same technique potentially has much broader uses, not only to find the colors of other source stars, but also in the timely recognition of high-magnification events and real-time analysis of anomalous events, which are both critical to the data-gathering stage of microlensing studies, as well as to the analysis of already-completed events. Finally, we have shown that the technique can be used derive otherwise unobtainable microlens parallaxes by aligning Earth-based and space-based lightcurves."
    },
    "0910/0910.0249_arXiv.txt": {
        "abstract": "Recent observations with Galaxy Evolution Explorer (GALEX) show strong unexpected UV excess in the spectrum of brightest cluster galaxies (BCGs).  It is believed that the excess UV signal is produced by old and evolved core-He burning stars, and the UV flux strength could be greatly enhanced if the progenitor stars have high value of He abundance.  In this work, we propose that sedimentation process can greatly enhance the He abundance in BCGs.  Our model predicts that the UV flux strength is stronger in more massive, low-redshift, and dynamically relaxed BCGs.  These predictions are testable with the current generation of GALEX+SDSS observations. ",
        "introduction": "\\label{sec:intro}  Brightest Cluster Galaxies (BCGs) are among the most massive galaxies in the Universe, forming and evolving at centers of galaxy clusters through galaxy mergers and accretion.  The baryonic component of present-day BCGs is composed of mainly old stellar populations \\citep[e.g.][]{DeLucia07, Collins09} with a fraction of younger stellar populations formed out of cooling gas in cluster cores \\citep[e.g.,][]{Peterson06}.  Properties of BCGs are therefore dependent on the mass accretion histories and the conversion of hot gas in clusters into stars.  Detailed studies of the BCGs properties have the potential to shed new insights into the physics of the most massive galaxies in the Universe.  The UV upturn is one of the mysterious phenomena observed in massive elliptical galaxies.  Since the first discovery \\citep{code_welch79}, it has been known for several decades that some giant elliptical galaxies exhibit strong unexpected bump in the Ultraviolet (UV) part of the spectrum \\citep[see][for a review]{OConnell99}.  Recent GALEX observations indicate that the UV flux excess is most significant in BCGs \\citep{ree_etal07}. The observed UV excess, also known as the UV upturn, is spatially smooth and extended \\citep[e.g.,][]{lee_etal05a,ree_etal07}.  The spatial distribution and the spectral shape of the UV upturn phenomena cannot be explained by recent star formation \\citep{brown_etal97}.  One of the leading models of the UV upturn suggests that the phenomena can be explained most naturally by old evolved stars, known as extreme horizontal-branch (EHB) stars.  EHB stars are hot, core helium-burning stars with extremely thin hydrogen envelopes ($M_{\\rm env} \\leq 0.05 M_\\sun$), which lost their envelopes via massive winds on the red giant branch in the single star evolution model \\citep[e.g.,][]{Dorman93} \\footnote{It has also been suggested that EHB stars can be produced in binary star systems, where the envelope of the core He-burning star can be tidally stripped by a companion star \\citep{Han07}.}.  High value of He abundance can make the formation of hot EHB stars more effective \\citep{Dorman95,Yi97}. This is because that HB stars with a higher helium abundance evolve faster on main sequence and red giant branch, can burn more hydrogen during the core helium-burning phase, and therefore can become UV-bright more easily.  In particular, a very large He abundance ($Y \\gtrsim 0.4$) can greatly boost the UV upturn strength \\citep{Dorman95}.  In this work, we point out that helium sedimentation process can greatly increase the helium abundance in BCGs, making them conducive to the formation of UV-bright EHB stars.  It has long been suggested that helium sedimentation occurs in the intracluster medium (ICM) of galaxy clusters \\citep{Abramopoulos1981On_the_equilibr,Gilfanov1984Intracluster,Qin2000BARYON-DISTRIBU,Chuzhoy2003Gravitational,Chuzhoy2004Element,Ettori2006Effects,Peng_Nagai09}. Under the influence of gravity, the heavier helium nuclei in the H-He dominated ICM accumulate at the center of massive galaxy clusters.  If the UV flux is produced mainly by EHB stars, the He sedimentation process can boost the UV flux produced by these stars.   Our sedimentation scenario predicts that the UV upturn phenomena should be most pronounced in high-mass (or $T_X$), low-redshift, and dynamically relaxed systems.  We show that these predictions are testable with the current generation of GALEX+SDSS observations. ",
        "conclusions": "\\label{sec:conclusion}  In this work, we propose that the He sedimentation can significantly enhance the UV signal observed in BCGs.  We show that the He sedimentation process can produce significant amount of He abundance ($\\Delta Y>0.25$) in central regions of hot, massive galaxy clusters. Stellar populations in BCGs may contain UV bright EHB stars when they are old.  The model predicts that the UV upturn phenomena increases with X-ray temperature or mass of galaxy clusters.  Our sedimentation model makes several unique predications that are testable with current GALEX+SDSS observations.  \\newcounter{bean} \\begin{list} {\\arabic{bean}.}{ \\usecounter{bean} \\setlength{\\parsep}{+0.03in} \\setlength{\\leftmargin}{+0.15in} \\setlength{\\rightmargin}{+0.15in}}  \\item{} The model predicts a strong correlation between the UV upturn strength of BCGs and the X-ray temperature of host galaxy clusters.  \\item{} The model also predicts that the UV upturn phenomena should be more pronounced in BCGs, compared to non-BCG elliptical galaxies.  \\item{} Some correlations are expected between the UV upturn strength of BCGs and the dynamical state of the BCGs and their host galaxy clusters.  Major mergers, for example, can destroy He-rich cluster cores, producing some scatter in the UV upturn among BCGs.  In this case, we expect a correlation between the size of UV upturn and dynamical state of clusters (e.g., measured using X-ray morphology or multiple X-ray peaks).  This is another expected observational signature of our model.  \\end{list}  Analyses of a large sample of galaxies from the GALEX+SDSS along with average temperatures of their host X-ray clusters can provide important test of this model and may shed new light on the origin of the UV upturn phenomena as well as the physics of the formation and evolution of massive galaxies in the universe."
    },
    "0910/0910.0908_arXiv.txt": {
        "abstract": "The discovery of bright gamma-ray emission coincident with supernova remnant (SNR) W51C is reported using the Large Area Telescope (LAT) on board the Fermi Gamma-ray Space Telescope. W51C is a middle-aged remnant ($\\sim 10^4$ yr) with intense radio synchrotron emission in its shell and known to be interacting with a molecular cloud. The gamma-ray emission is spatially extended, broadly consistent with the radio and X-ray extent of SNR W51C. The energy spectrum in the 0.2--50 GeV band exhibits steepening toward high energies. The luminosity is greater than $1\\times 10^{36}\\ \\rm erg\\ s^{-1}$ given the distance constraint of $D>5.5$ kpc, which makes this object one of the most luminous gamma-ray sources in our Galaxy. The observed  gamma-rays  can be explained reasonably by a combination of efficient acceleration of nuclear cosmic rays at supernova shocks and shock-cloud interactions. The decay of neutral $\\pi$-mesons produced in hadronic collisions provides a plausible explanation for the gamma-ray emission. The product of the average gas density and the total energy content of the accelerated protons amounts to $\\bar{n}_{\\rm H}W_p \\simeq 5\\times 10^{51}\\ (D/6\\ {\\rm kpc})^2\\ \\rm erg\\ cm^{-3}$. Electron density constraints from the radio and X-ray bands render it difficult to explain the LAT signal as due to inverse Compton scattering. The \\emph{Fermi} LAT source coincident with SNR W51C sheds new light on the origin of Galactic cosmic rays. ",
        "introduction": "The origin of cosmic rays  remains unsolved. The idea that supernovae produce cosmic rays  \\citep[e.g.,][]{Hayakawa56} has been developed both in theoretical and observational aspects, so that galactic cosmic rays are widely considered to be accelerated in the expanding shocks of supernova remnants (SNRs) through diffusive shock acceleration process \\citep[see e.g.,][]{BE87}. This conjecture has been strengthened by recent observations of young SNRs in synchrotron X-rays and very-high-energy (VHE) gamma-rays \\citep[see e.g.,][]{Reynolds08,Aha08}. High efficiency for converting the kinetic energy released by supernova explosions into the energy of relativistic protons and nuclei is a key requirement to explain the galactic cosmic rays; it is yet to be confirmed observationally.  Enhanced $\\pi^0$-decay emission expected from SNRs that are interacting with  molecular clouds could provide direct evidence of the nuclear component of cosmic rays being accelerated at supernova shocks \\citep{ADV94}. To identify the  $\\pi^0$-decay gamma-rays as evidence of the nuclear cosmic rays, observations in the GeV domain are crucial because the  $\\pi^0$-decay spectrum has a characteristic  bump around 70 MeV. Previous  measurements of GeV gamma-rays with EGRET onboard the Compton Gamma-Ray Observatory found some gamma-ray sources near radio-bright SNRs \\citep{Esp96}. However, the possible origins of the EGRET sources, namely SNR shell contributions or pulsar associations, remained unclear, mainly due to poor localization.  The advent of the Large Area Telescope (LAT) onboard the \\emph{Fermi} Gamma-ray Space Telescope provides a new opportunity to study the gamma-ray emission from SNRs at GeV energies. An initial source list based on the first three months of \\emph{Fermi} observations \\citep{BSL} includes 0FGL~J1923.0$+$1411, which is spatially coincident with SNR W51C. Even with the three months data, the detection was at $\\sim 23\\sigma$ level. There are no EGRET counterpart(s) to the LAT source \\citep{EGRET3}; this is the first SNR discovered by \\emph{Fermi} as confirmed in this paper.  W51C (G49.2$-$0.7) is a radio-bright SNR at a distance of $D \\simeq 6$ kpc with an estimated age of $\\sim 3\\times 10^4$ yrs \\citep{Koo95}. It has an elliptical shape with an extent of $50\\arcmin \\times 38\\arcmin$. The radio continuum map exhibits  thick shell-like structures \\citep{330MHz}. The bulk of the X-ray emission comes from thermal plasma with a temperature of $kT\\sim 0.3\\ {\\rm keV}$ \\citep{Koo05}. A molecular cloud--shock interaction is known, as evidenced by observations of shocked atomic and molecular gases \\citep{KM97a,KM97b}. Most recently, the HESS collaboration has announced the detection of  extended VHE gamma-ray emission coincident with W51C  \\citep{HESSW51C}. Also, the Milagro collaboration has reported a possible excess of multi-TeV gamma-rays in this direction \\citep{Milagro09}.  In this paper we report the analysis results for the LAT source coincident with SNR W51C, using data accumulated over the first year of \\emph{Fermi}'s operation. The \\emph{Fermi} observations of SNR W51C  permit a refined study of cosmic-ray acceleration. Specifically, the LAT data suggest that $\\pi^0$-decay emission is the dominant contribution to the gamma-ray signal. ",
        "conclusions": "\\label{sec:discuss}  The extended gamma-ray emission positionally coincident with SNR W51C has been studied using the \\emph{Fermi} LAT. The gamma-ray spectrum presented in Figure~\\ref{fig:SED} is not fit by a simple power law, exhibiting a remarkable steepening. Here we discuss the origin of the extended emission and the underlying particle spectra that give rise to the observed  spectrum of photons.   The expanding shock waves driven by the supernova explosion are expected to be the sites of the acceleration of multi-GeV particles. To phenomenologically interpret the spectral curvature in the LAT+TeV bands, a broken power-law is adopted for the momentum distribution of the radiating electrons/protons: \\begin{equation} \\label{eq:N} N_{e,p}(p) = a_{e,p} \\left(\\frac{p}{p_0}\\right)^{-s} \\left(1+  \\left(\\frac{p}{p_{\\rm br}}\\right)^{2} \\right)^{-\\Delta s/2}, \\end{equation} where $p_0 = 1\\ {\\rm GeV}\\, c^{-1}$. For simplicity, the indices and the break momentum are assumed to be identical for both accelerated protons and electrons. As we argue below, the break may reflect the character of magnetohydrodynamic turbulence. To account for  the radio synchrotron index $\\alpha \\simeq 0.26$ \\citep{MK94}, we adopt $s=1.5$, though a steeper index (say, $s = 1.7$) could be reconciled  with the radio observations within the uncertainty. The energetic particles are assumed to be uniformly distributed in the volume of $V=(4\\pi /3) R_{\\rm eff}^3$ with an effective radius $R_{\\rm eff}=30$ pc. The age and radius imply a shock velocity of $v_{\\rm sh} \\sim 400\\ {\\rm km\\ s^{-1}}$ and $E_{\\rm SN}/n_0 \\sim 1.6\\times 10^{52}\\ {\\rm erg\\ cm^{3}}$, where $E_{\\rm SN}$ and $n_0$ represent the explosion kinetic energy and the interstellar density into which the main blast wave is propagating, respectively. The X-ray data suggest $n_0 \\sim 0.3\\ {\\rm cm^{-3}}$. The radio images indicate that radio-emitting electrons are smoothly distributed in a thick shell. The model of \\citet{KM97a} suggests the presence of a molecular cloud of $\\sim 1\\times 10^4\\ M_{\\odot}$ engulfed by the  blast wave. The engulfed cloud can act as target material for relativistic particles. The total (atomic and molecular) hydrogen mass contained in the  volume is denoted by $M_{\\rm H}= \\bar{n}_{\\rm H} m_p V$. Note that $\\bar{n}_{\\rm H} \\gg n_0$ can be expected.  Given the interaction with a molecular cloud, we first attribute the observed gamma-rays to the decay of $\\pi^0$ mesons produced in inelastic collisions between accelerated protons (and nuclei) and target gas (Fig.~\\ref{fig:SED}). The gamma-ray spectrum of $\\pi^0$-decay is calculated based on \\citet{Kamae06} using a scaling factor of 1.85 for helium and heavy nuclei \\citep{Mori09}. Note that the scaling factor assumes the local interstellar abundance for target material and the observed cosmic-ray composition. Contributions from bremsstrahlung and Inverse-Compton (IC) scattering by accelerated electrons are also shown in Fig.~\\ref{fig:SED}. Electron-ion and electron-electron bremsstrahlung spectra are computed as in \\citet{Baring99}. The  interstellar radiation field for IC (see Table~\\ref{tbl:model})  is comprised of two diluted blackbody components (infrared and optical) and the CMB. The infrared and optical components are adjusted to reproduce the interstellar radiation field in the GALPROP code \\citep{Porter08}. Cooling effects due to ionization and synchrotron (or IC) losses, which introduce cooling breaks in particle spectra in addition to $p_{\\rm br}$, are taken into account assuming constant particle injection over a period $T_0 \\sim {3\\times 10^4}$ yr. The synchrotron cooling becomes important in the TeV band for leptonic models.  Figure \\ref{fig:multiband} (a) shows the radio+gamma-ray spectrum together with the radiation model that uses the parameters in Table~\\ref{tbl:model}. We  adopt here $M_{\\rm H} = 2.8\\times 10^4\\ M_{\\odot}$ ($\\bar{n}_{\\rm H}=10\\ {\\rm cm^{-3}}$), which is somewhat larger than the value quoted above. The total energy of the high-energy protons amounts to $W_p = 5.2\\times 10^{50}\\ {\\rm erg}$, which is inversely proportional to $M_{\\rm H}$, but insensitive to other parameters. The large luminosity of the LAT source can be explained naturally by a large number of accelerated protons and dense environments as expected for the case of W51C.  The break feature in the proton spectrum, which is introduced on phenomenological grounds,  is not expected in a shock acceleration zone if acceleration proceeds close to the Bohm limit \\citep[e.g.,][]{Baring99}. Protons can be accelerated up to $\\sim 200$ TeV  via diffusive shock acceleration with a power-law (or even slightly concave) momentum distribution. A possible explanation for the falling proton spectrum above $p_{\\rm br} = 10\\mbox{--}30\\ {\\rm GeV}\\, c^{-1}$ (which depends on the choice of other parameters) could be the effects  of damping of magnetohydrodynamic turbulence due to ion-neutral collisions. According to \\citet{PZ03}, the maximum attainable momentum in the presence of wave dissipation by ion-neutral collisions is estimated as $p_{\\rm max}(T_0) \\sim 30\\, (n/1\\, {\\rm cm^{-3}})^{-1/2}\\ {\\rm GeV}\\, c^{-1}$, using the SNR parameters described above. Here $n$ is the ambient neutral gas density. The shock precursor cannot confine accelerated protons with $p>p_{\\rm max}(T_0)$ even if they were accelerated earlier. Therefore, energy-dependent escape of accelerated particles from the remnant \\citep[e.g.,][]{AA96} may account for the steepening of the proton distribution. In any case, a break is needed to fit the observed SED.  In contrast, leptonic models for the GeV gamma-ray production face difficulties (see Figure \\ref{fig:multiband}b--c and Table~\\ref{tbl:model}). While the chosen model parameters are not unique in their ability to fit the broadband spectrum, they are representative; no reasonable choices for them can easily account for the radio+gamma-ray data. Leptonic scenarios require $a_e/a_p$ far in excess of cosmic-ray abundance ratios. It is difficult to reconcile the bremsstrahlung-dominated model with the observed radio synchrotron spectrum  (Fig.~\\ref{fig:multiband}b). Also, the IC-dominated model (Fig.~\\ref{fig:multiband}c) requires an unrealistically large  energy content in radiating electrons, $W_e \\simeq  1\\times 10^{51}\\ {\\rm erg}$, a low magnetic field of $B \\simeq 2\\ \\mu{\\rm G}$ (to  suppress the radio synchrotron emission), and a low density of $\\bar{n}_{\\rm H} < 0.1\\ {\\rm cm^{-3}}$ (to reduce the bremsstrahlung component). Such a low density is problematic because even X-ray-emitting gas has a density of $\\sim 1\\ {\\rm cm^{-3}}$ \\citep{Koo02}. It is also difficult to account for the LAT signal by an IC process in a possible pulsar wind nebula seen inside W51C,  \\CXOJ\\ \\citep{Koo05}, which has only a few parsec extent and much weaker radio emission compared with the SNR shell.  The most plausible explanation for the LAT source is therefore offered by $\\pi^0$-decay gamma-rays in a dense environment, spawned by efficient acceleration of protons and nuclei taking place or having taken place at the shocked shell of SNR W51C. The discovery of the GeV-scale gamma-rays presented in this paper raises  the intriguing possibility that {\\it Fermi} has discerned accelerated ions in a middle-aged remnant that is interacting with dense clouds."
    },
    "0910/0910.2416_arXiv.txt": {
        "abstract": "We present an analysis of the neutral hydrogen and stellar populations of elliptical galaxies in the \\cite{2009arXiv0908.1382T} sample. Our aim is to test their conclusion that the continuing assembly of these galaxies at $z$$\\sim$0 is essentially gas-free and not accompanied by significant star formation. In order to do so, we make use of \\hi\\ data and line-strength indices available in the literature. We look for direct and indirect evidence of the presence of cold gas during the recent assembly of these objects and analyse its relation to galaxy morphological fine structure.  We find that $\\geq$25\\% of ellipticals contain \\hi\\ at the level of \\mhi$>$10$^8$\\msun, and that \\mhi\\ is of the order of a few percent of the total stellar mass. Available data are insufficient to establish whether galaxies with a disturbed stellar morphology are more likely to contain \\hi. However, \\hi\\ interferometry reveals very disturbed gas morphology/kinematics in all but one of the detected systems, confirming the continuing assembly of many ellipticals but also showing that this is not necessarily gas-free. We also find that all very disturbed ellipticals have a single-stellar-population-equivalent age $<$4 Gyr. We interpret this as evidence that $\\sim$0.5-5\\% of their stellar mass is contained in a young population formed during the past $\\sim$1 Gyr. Overall, a large fraction of ellipticals seem to have continued their assembly over the past few Gyr in the presence of a mass of cold gas of the order of 10\\% of the galaxy stellar mass. This material is now observable as neutral hydrogen and young stars. ",
        "introduction": "The role of dissipationless galaxy merging in the assembly of early-type galaxies is currently under debate. On the one hand, it appears to be necessary to produce the boxy, slowly-rotating galaxies which populate the high-mass end of the red sequence \\citep[e.g., ][and references therein]{1997AJ....114.1771F,2009ApJS..182..216K}. On the other hand, dissipative processes are required to explain a number of structural, kinematical and stellar-population properties of intermediate- and low-mass early-type galaxies \\citep[e.g., ][and refernces therein]{1992ApJ...399..462B,2007MNRAS.379..401E,2009ApJS..181..135H}. Furthermore, the tight scaling relations observed in the local Universe place a rather conservative upper limit on the fraction of stellar mass assembled via dissipationless merging \\citep[e.g., ][]{2003MNRAS.342..501N,2009arXiv0908.1621N}. Finally, cold gas is actually observed in a large fraction of early-type galaxies in the nearby Universe \\citep[e.g., ][]{2006MNRAS.371..157M,2007MNRAS.377.1795C}.  Based on deep imaging of a volume-limited sample of nearby elliptical galaxies with M$_B \\leq-20$, and following earlier work by, e.g., \\cite{1983ApJ...274..534M} and \\cite{1992AJ....104.1039S}, \\citet[hereafter T09]{2009arXiv0908.1382T} find signatures of recent dynamically-violent assembly in a remarkably large fraction of objects, $\\sim$75\\%. Analysing galaxies' $B-V$ colour, they argue that ellipticals outside cluster environments continue to grow at $z$=0 mostly through gas-free accretions. If correct, this conclusion would have implications for a number of properties of local ellipticals. This follows from a large number of numerical studies showing that the presence of even a modest amount of gas ($\\sim$10-20\\% of the baryonic mass)  in a binary merger can change the stellar mass-distribution and orbits in the early-type remnant relative to the gas-free case \\citep[e.g., ][]{2005MNRAS.360.1185J,2007MNRAS.376..997J,2009ApJS..181..135H}. Observable quantities such as ellipticity, isophote disciness/boxiness, anisotropy and light profile are affected by the gas content of the progenitor galaxies.  Although most of these numerical simulations focus on binary major or minor mergers, early-type galaxies in a $\\Lambda$CDM Universe form following a sequence of many interactions with widely varying mass ratio and degree of dissipation. In other words, galaxies are much more likely to grow via continuous assembly of smaller units of varying gas content rather than a few major events, in a way strongly dependent on their environment \\citep[e.g., ][]{2006MNRAS.366..499D,2007ApJ...658..710N,2009MNRAS.396.1972P}. Within this framework, as in the binary merger case, a galaxy formed following a gas-free path (e.g., one residing in the centre of a massive cluster, where little cold gas can penetrate) should be significantly different from one that experienced dissipation \\citep[e.g.,][]{2007A&A...476.1179B}. It is therefore interesting to estimate the mass of cold gas present during the continuing assembly of ellipticals over the last few Gyr and test T09 conclusion that this is indeed negligible.  The approach taken in this Letter is to look for neutral hydrogen and young stellar populations in galaxies in the T09 sample. Both properties could be regarded as marginally important with respect to the broad cosmological picture of galaxy evolution. After all, work over the past decades shows that young stars and cold gas amount to just a few percent of the total galaxy stellar mass of $z$=0 ellipticals. However, as explained above, even modest amounts of gas can have important effects on fundamental properties of these objects. We would like to clarify what kind of data analysis is necessary in order to detect this gas. Work along this line has been carried out by a number of authors during the past few years. We summarise it together with our results in the discussion section. ",
        "conclusions": "Neutral hydrogen observations and line-strength index analysis of the stellar populations of ellipticals in the T09 sample reveal that a large fraction of these galaxies continued their assembly during the past few Gyr in the presence of a mass of cold gas of \\emph{at least} a few percent of the total galaxy stellar mass. This gas is now detected in the form of \\hi\\ and through young stellar sub-populations.  \\hi\\ interferometry  reveals that in nearly all cases the detected gas is morphologically/kinematically disturbed and mostly found in tails/filaments, sometimes clearly associated to the optical morphological disturbances. This, and the presence of young stars, confirm T09 conclusion that the assembly of  elliptical galaxies continues to $z$=0. However, it does not agree with their conclusion that  this assembly is essentially gas-free. In fact, \\hi\\ observations could be regarded as yet one more tool to reveal the on-going assembly of these systems, complementary to deep optical imaging \\citep[see][]{2006MNRAS.371..157M,2007A&A...465..787O,2009A&A...498..407G}.  Our estimate of the mass of cold gas recently accreted by ellipticals is in qualitative agreement with the mass accretion rate derived by T09 themselves. They estimate that nearby ellipticals are still growing at $z$=0 at a rate of 20\\%-in-mass per Gyr. Evidence of this on-going assembly is particularly strong in the field and in galaxy groups, where most small galaxies around a given elliptical are gas rich. It is therefore reasonable to expect that such a conspicuous mass assembly is accompanied by an accretion of cold gas of \\emph{at least} a few percent of the total galaxy stellar mass.  Despite this high incidence of disturbed \\hi\\ systems, elliptical galaxies with very regular gas distributions do exist \\citep[e.g., NGC~2974; more examples can be found in ][]{2006MNRAS.371..157M,2007A&A...465..787O}. These systems could be evolved versions of the galaxies that show disturbed \\hi\\ at $z$=0, i.e., galaxies for which gas accretion was important many Gyr ago and whose recent history has been quiet enough for this gas to settle.  Another interesting point is that while we find that optically disturbed ellipticals are systematically younger than relaxed objects, and interpret this as evidence of their continuing, gas-rich assembly, not all ellipticals with \\hi\\ are young (e.g., NGC~5846; other cases are known outside the T09 sample, e.g., NGC~4278 in \\citealt{2006MNRAS.371..157M}). Clearly, the dynamics of the accretion plays a very important role in determining the effect of gas on structure and (distribution of) stellar population of the accreting galaxy. Understanding where the young stars and cold gas are, and what their kinematics is, is as important as estimating their mass.  Our results agree with previous work on this subject. \\cite{2009arXiv0908.2548S} analyse the stellar population of high-$z$ early-type galaxies selected by \\cite{2005AJ....130.2647V} to be ``dry''-merger remnants. They find that highly disturbed objects are on average younger than relaxed ones, and that their line-strength indices are consistent with the presence of a few-percent-in-mass of young stars. This is in agreement with early results from \\cite{1990ApJ...364L..33S} who find the \\hb\\ index to be systematically higher for ellipticals with morphological fine structure at any given absolute magnitude. Finally, \\cite{2007AJ....134.1118D} look at the \\hi\\ properties of a sample of local galaxies selected in the same way as the \\cite{2005AJ....130.2647V} objects. They find that many of these systems host a significant mass of \\hi\\ whose disturbed morphology and kinematics make them consistent with having recently accreted cold gas.  Following the conclusions of a large number of numerical studies, the direct or indirect detection of gas present during the assembly of elliptical galaxies may have important implications for many structural properties of these objects. We have shown that using only observables such as colours measured over large apertures one risks to neglect the evidence of this gas because of the lower sensitivity to the presence of a modest, but nevertheless significant, mass of cold gas or young-stars."
    },
    "0910/0910.5236_arXiv.txt": {
        "abstract": "We investigate the effect of an interaction between dark energy and dark matter upon the dynamics of galaxy clusters.  This effect is computed through the Layser-Irvine equation, which describes how an astrophysical system reaches virial equilibrium and was modified to include the dark interactions. Using observational data from almost 100 purportedly relaxed galaxy clusters we put constraints on the strength of the couplings in the dark sector. We compare our results with those from other observations and find that a positive (in the sense of energy flow from dark energy to dark matter) non vanishing interaction is consistent with the data within several standard deviations. ",
        "introduction": "One of the most surprising results of physics and cosmology in the last ten years is the cosmological accelerated expansion, which has been proved beyond reasonable doubt by observations \\cite{cosmoaccel,wmap,bao}. The simplest explanation of such observations is a cosmological constant, which is, however, off theoretical computations by 120 orders of magnitude, a result that calls for urgent explanations. Another possibility is to argue that the Universe contains a strange dynamical component with negative pressure, dark energy, which is responsible for more than three quarters of its energy content \\cite{darkenergy}.   Interactions of dark energy or dark matter with baryonic matter and radiation must be either inexistent or negligible in order to comply with stringent observations of the visible sector.  However, in a field-theoretical approach, dark matter is some particle of a unification scheme, and dark energy should also be described by a field, which means that some level of interaction between these different sectors is basically mandatory. In the literature there is a large body of work dealing with such a possibility, e.g., \\cite{amendola}-\\cite{[14]}. It is interesting to notice that a coupling between dark energy and dark matter can also serve to alleviate the coincidence problem \\cite{amendola,[11]}. The value of the coupling and holographic arguments also allow for the crossing of the phantom barrier which separates models with equations of state $w \\equiv p/\\rho > -1$ from models with $w < -1$ \\cite{[12]} -- see also \\cite{[14],[15]}. Further consequences of the interaction have been studied, as {\\it e.g.} its effect on the lowest multipoles of the cosmic microwave background (CMB) spectrum \\cite{[13],[16]}.  The strength of the coupling is presumably as small as the fine structure constant \\cite{[13],[17]}. Comparison with supernova data and CMB and large-scale structure have also been analysed \\cite{[18]}. Nevertheless, the observational limits on the strength of such an interaction remain weak \\cite{[19]}.   In addition to these constraints, dynamical dark energy (i.e., a time- and space-dependent field) has an impact on the large-scale structure due to its fluctuations. In that case, dark energy affects the process of structure formation by means of its density fluctuations, both in the linear \\cite{[11]}, \\cite{[20]}-\\cite{[23]} and the non-linear \\cite{[24],[25]} regimes, as the growth of dark matter perturbations can be affected \\cite{[13],[14],[26]}.   However, most constraints on dark interactions concern the asymptotic  behaviour of the dark sector, both in time and in space. Local or nearby checks are still weak. This started to change with the detailed analysis of galaxy clusters and their internal structure. It was suggested that the dynamical equilibrium of collapsed structures could be affected by the coupling of dark energy to dark matter, which as a result could affect the averaged energy distributions predicted by the virial theorem -- in fact, by its relativistic counterpart, the Layser-Irvine equation \\cite{[28]}. This was first proposed, and analyzed for the Abell cluster A586, by \\cite{[27]}.  The basic idea is that the virial theorem is distorted by the mass non-conservation generated by the coupling of dark matter with dark energy.   In a previous paper \\cite{ababrasw} we showed how the Layser-Irvine equation, describing the evolution towards virialization, is changed by the presence of the coupling. We showed that this violation leads to a systematic bias in the estimation of virial masses of clusters when the usual virial conditions are employed. By comparing weak lensing and X-ray mass-observables on the one hand, and virial masses on the other hand, we were able to cross-check not only the strength of the coupling, but also whether it was indeed a feature of the virial mass estimate, and not a more complicated set of biases between all these mass-observable relations. Our main result was that a single parameter could explain all the bias between virial mass and the other estimates (weak lensing and X-ray) to 95\\% confidence level (C.L.), or, equivalently, two standard deviations.  However, the amount of data (about thirty independent useful observations) did not allow a detailed numerical analysis, resulting in relatively weak constraints on the coupling. In this paper we will not only reanalyze that data, but will also include a new set of data for almost a hundred X-ray observations of galaxy clusters, and compare the presumed mass from those observations with the virial masses for those clusters.  Notice that even though the uncertainties associated with any individual galaxy cluster are very large, by comparing the naive virial masses of a large sample of clusters with their masses estimated by X-ray and by weak lensing data, we are able to constrain the bias between them and to impose much tighter limits on the strength of the coupling than has been achieved before. ",
        "conclusions": "We are fully aware that our lack of precise knowledge about the errors arising from systematics could weaken our conclusions. We recall here that our main work hypothesis is that the differences between the various kinds of masses estimates come solely from the fact that they do or do not take into account the interaction parameter. Obviously, this assumption is not absolutely true, but represents an approximation we use to replace the need to analyze any particular systematic effects. Also, we have used in this work the best available data sets well-suited for our purposes. Nevertheless, we still believe that our analysis, combined with previous results, signals in a coherent manner towards the possibility of the dark energy-dark matter interaction.  The scenario in which dark energy and dark matter are in interaction seems to be a strong physical possibility.  The interaction intensity is a small constant, possibly larger than the fine structure constant, in agreement with \\cite{[17]}. The results are all consistent among themselves, and we obtain the interesting conclusion that the interaction constant is far from zero by several (possibly three) standard deviations, which is a rather intriguing conclusion. The positive value is also in agreement with the thermodynamical arguments of \\cite{wangpavon}, strengthening our conclusions. A few words should be said about the possible influence of the interaction parameter on the cosmic history (in particular, in the structure formation process). The behavior of dark matter density as a function of time can be rather different from the non interacting case depending on the kind of interaction considered. We do not think that the proposed interaction should be the same along the whole history since decoupling, when a more robust suggestion has to be made. Since usually the matter-radiation equality is supposed to occur, in the standard model (e.g. in the non dark sector case) as well, the order of magnitude of the cosmological time parameter should not get strongly modified. (Notice that in the past, with interaction, dark matter density was smaller compared to the noninteracting case.  Thus, the matter density is bounded from below by the baryonic matter density.)"
    },
    "0910/0910.1286_arXiv.txt": {
        "abstract": "The 6307\\AA\\ emission line in the spectrum of \\ec\\ \\citep{Martin06a} is blue-shifted [\\altion{S}{iii}] \\lm6313 emission originating from  the outer wind structures of the massive binary system. We realized the identification while analyzing multiple forbidden emission lines not normally seen in the spectra of massive stars. The high spatial and moderate spectral resolutions of \\hst/STIS resolve forbidden lines of Fe$^+$, N$^+$, Fe$^{+2}$, S$^{+2}$, Ne$^{+2}$ and Ar$^{+2}$ into spatially and velocity-resolved  rope-like features originating from collisionally-excited ions photo-ionized by UV  photons or collisions. While the [\\altion{Fe}{ii}] emission extends across  a velocity range of $\\pm$500~\\kms\\ out to 0\\farcs7, more highly ionized forbidden emissions ( [\\altion{N}{ii}], [\\altion{Fe}{iii}], [\\altion{S}{iii}], [\\altion{Ar}{iii}], and [\\altion{Ne}{iii}]) range in velocity from $-$500 to $+$200 \\kms, but spatially extend outward to only  0\\farcs4. The [\\altion{Fe}{ii}] defines the outer regions of the massive primary wind. The [\\altion{N}{ii}], [\\altion{Fe}{iii}] emission define the  the outer wind interaction regions directly photo-ionized by far-UV radiation. Variations in emission of [\\altion{S}{iii}] \\lm\\lm 9533, 9071 and 6313 suggest density ranges of 10$^6$ - 10$^{10}$~\\den\\ for electron temperatures ranging from 8,000 to 13,000\\degr\\ K. Mapping the temporal changes of the emission structure at critical phases of the 5.54-year period  will provide important diagnostics of the interacting winds. ",
        "introduction": "\\citet{Martin06a} found  a variable emission line  centered at 6307\\AA\\ in multiple spectra of \\ec, recorded by \\hst/STIS and by \\vlt/UVES. They were unable to identify the origin of the emission line, but demonstrated that the line was present across the high state (defined by presence of forbidden lines of doubly-ionized elements, see \\cite{Damineli08a} and references therein) and disappeared during the low state. \\\\ \\indent Recently \\citet{Gull09a}, using the same spectra, focused on the spatially-extended forbidden line emission both from high-ionization (herein defined as $>$14 eV) and low-ionization (8-13 eV). We found that the forbidden emission originated from 1)  the Weigelt condensations \\citep{Weigelt86},  as narrow lines  centered on $-$43 \\kms, 2)  the boundaries of the primary wind (\\ec~A) as rope-like [\\altion{Fe}{ii}], photo-ionized by mid-UV and collisionally excited at densities around N$_c$=10$^7$\\den, and 3) the wind interaction region, as high-ionization emission  from [\\altion{N}{ii}], [\\altion{Fe}{iii}], [\\altion{Ar}{iii}], [\\altion{Ne}{iii}] and [\\altion{S}{iii}], photo-ionized by far-UV and collisionally excited for densities, n$_c$ ranging from 10$^5$ to 10$^8$\\den.  The high-ionization emission lines are present both for the Weigelt condensations and the wind interaction region during the 5 year high state, but every 5.5 years, disappear during the low state. X-ray models \\citep{Pittard02, Parkin09a} place the binary periastron event near the onset of the low state with a three to six month recovery.\\\\ \\indent We applied a 3D SPH (smoothed particle hydrodynamics) model \\citep{Okazaki08a} extended out to 1700 AU (0\\farcs67) to match the spatial structure seen in these forbidden lines and realized that the bulk of the high-ionization emission structure originated from the wind interaction region in the outer portion of the massive wind structure. We  found that portions of the wind interaction structure,  moving ballistically outward, are directly illuminated by the far-UV radiation of the hot secondary, \\ec~B, leading to  highly-ionized, collisionally excited gas and hence the high-ionization extended emission.\\\\ \\indent Further examination of the \\hst/STIS longslit spectra showed that the previously unidentified emission at 6307\\AA\\ is  blue-shifted [\\altion{S}{iii}] $\\lambda$6313 emission from the interacting wind structure. The evidence is subtle, but convincing. We summarize the observations in Section 2. A description of how the spectro-images were produced is  in Section 3. The forbidden emission structures are described in Section 4. Discussion in Section 5 provides insight on the ionization and excitation leading to [\\altion{S}{iii}] emission and  the potential for monitoring changes with orbital phase, including mapping temperature with density dependence. We conclude with a summary in Section 6. Throughout this paper, all wavelengths are  in vacuum, the velocities are heliocentric, directions are compass points (N=north, NNW= north by northwest and the phase of the binary orbit is referenced to the X-ray minimum beginning at 1997.9604 \\citep{Corcoran05a}. \\vfill\\eject ",
        "conclusions": "We have presented conclusive evidence that the emission line at 6307\\AA, noted in the spectra of \\ec\\ by \\citet{Martin06a} is blue-shifted emission of [\\altion{S}{iii}] \\lm6313 originating from the distorted paraboloidal, interaction region located between the massive binary members. While the massive primary, \\ec~A, provides the dominant wind ejecta, 10$^{-3}$ \\Msun/y at 500 \\kms, the hotter secondary provides a less massive, faster wind, 10$^{-5}$ \\Msun/y at 3000 \\kms, and far-UV photons that ionize iron, neon, argon and sulfur to doubly ionized states. Thermal collisions, mid-UV photons and possible charge exchange excite the doubly-ionized species to upper states with forbidden transitions leading to forbidden line emission in regions with densities close to n$_c$. Specifically, [\\altion{S}{iii}] \\lm\\lm9533, 9071 and 6313 have extended spatial structures. The intensity ratio (Flux (\\lm9533) + Flux (\\lm9071)/Flux(\\lm6313) leads to density estimates ranging from 10$^7$ - 10$^8$ \\den on the scale of 0\\farcs1, the limit of \\hst/STIS spatial capabilities. Mapping in these and other doubly-ionized lines will provide powerful measures for models of wind interactions using various 3-D hydrodynamical codes."
    },
    "0910/0910.3839.txt": {
        "abstract": "Recent observation  of microlensing technique reveals  two giant planets at  2.3 AU and 4.6 AU around the star OGLE-06-109L. The eccentricity of the outer planet ($e_c$) is estimated to be 0.11$_{-0.04}^{+0.17}$, comparable to  that of Saturn (0.01-0.09). %with unknown eccentricity for the inner planet . The similarities between the OGLE-06-109L system and the solar system indicate that they may have passed through  similar histories during their formation stage.  In this paper we investigate the dynamics and formation of the orbital architecture in the OGLE-06-109L system. For the present two planets with their nominal locations, the secular motions are stable as long as their eccentricities ($e_b,~e_c$) fulfill $e_b^2+e_c^2 \\le 0.3^2$. Earth-size bodies might be formed and are stable in the habitable zone (0.25AU-0.36AU) of the system. Three possible scenarios may be accounted for formation of $e_b$ and $e_c$: (i) convergent migration of two planets and the 3:1 MMR trapping; (ii) planetary scattering; (iii) divergent migration and the 3:1 MMR crossing. As we showed that the probability  for the two giant planets in 3:1 MMR is low ($\\sim 3\\%$),  scenario (i) is less likely. According to models (ii) and (iii), the final eccentricity of inner planet ($e_b$) may oscillate between [0-0.06], comparable to that  of Jupiter (0.03-0.06). An inspection of $e_b$, $e_c$'s secular motion may be helpful to understand which model  is really responsible for the eccentricity formation. ",
        "introduction": "To date, more than 370 exoplanets are detected mainly by Doppler radial velocity measurements(\\footnote{http://exoplanet.eu/}), among them  9 planets are revealed by gravitational microlensing (Udalski et al. 2005; Beaulieu et al. 2006; Bennett et al. 2006, 2008; Gould et al. 2006; Gaudi et al. 2008; Dong et al. 2009; Janczak et al. 2009). The use of microlensing   technique to  planet-searching is based on the idea of general relativity  that light passing through a mass should bend as if it passes a lens (Mao \\& Paczynski 1991). The major advantage of the technique is that it favors to detect planets in moderate distance to the host star ($\\sim$ 3 AU), which is complimentary to the radial velocity measurements. Among the planets detected by microlensing, OGLE-06-109L system is the first one with observed multiple planets.  It is 1490 pc away from the Sun, with a star of $\\sim 0.5~ M_\\odot$ (solar mass) and  two planets of $0.71~M_J$ (Jupiter mass) and 0.27 $M_J$ in the orbits of 2.3 AU and 4.6 AU, respectively (Gaudi et al. 2008).  Table 1 lists their nominal orbital elements.  Several features of the system make it an analogy of the solar system.  (i)  The two planets have positions and masses similar to those of Jupiter and Saturn up to  scale changes. (ii) The fitted eccentricity of  OGLE-06-109L c  is modest (0.11$_{-0.04}^{+0.17}$), comparable with those of Jupiter (0.03-0.06) and Saturn (0.01-0.09) during their secular evolution, although the eccentricity of  OGLE-06-109L b  is unknown. (iii) While Jupiter and Saturn's orbits are closing to $5:2$ mean motion resonance (MMR), OGLE-06-109L b and c may be close to $3:1$ MMR. (iv) All these four giant planets are  located well outside the snow lines of their systems ($\\sim 0.68$ AU for OGLE-06-109L and $\\sim 2.7$ AU for  solar system),   indicating that, unlike the most observed extra-solar systems with hot Jupiters, the migration of these four giant planets were not so efficient. (v) The habitable zone of the OGLE-06-109L system,  [0.25AU-0.36AU], may have stable orbits.  Considering most of the observed exoplanets have close-in orbits with an average eccentricity $\\sim 0.2$ (Udry \\& Santos  2007),  the  investigation of formation scenario in  OGLE-06-109L system may bridge the gap between the solar system and the most observed  exoplanet systems. Another aim of the paper is to predict the eccentricity of OGLE-06-109L b. Due to the significant uncertainties  in  orbital determination of microlensing technique ($\\sim 10\\%$,  S.  Mao, private communication), the orbital parameters of the two detected planets in OGLE-06-109L system  are poorly known. A precise determination  by other means like radial velocity is still impossible due to the great distance ($\\sim 1490$ pc) between OGLE-06-109L and the Sun. So the investigation of their formation and dynamics is  helpful to reveal their orbital parameters, especially their eccentricities.  The mechanism of eccentricity excitation is not fully understood, especially in single planet systems. During the early stage of planet formation, disk-planet interaction tends  to damp the eccentricity of the planet. According to the  linear theory, the planet exerts torques on a dynamically cold disk ($c\\ll r\\Omega $, where $c,~r,~\\Omega$ are the sound speed of gas, the orbital radius and angular velocity of planet motion, respectively) mainly at the Lindblad (LR) and corotation (CR) resonances (Goldreich \\& Tremaine, 1979, 1980; Ward 1988). Only non co-orbital LRs can excite planetary eccentricity, while CRs and co-orbital LRs  damp the eccentricity. In the linear regime, torques from the latter dominate the evolution, so the planetary eccentricity  is damped, unless  the planet is massive enough to clear the local gas disk (Lin \\& Papaloizou 1993). Two-dimensional hydrodynamical simulations show that the critical mass ($M_{\\rm crit}$) of the planet above which its eccentricity will be excited under disk tide  is $\\sim 20~ M_J $, and $M_{\\rm crit}$ might be reduced into the range of the  observed extrasolar planets at a very low disc viscosity (Papaloizou et al. 2001). During the later stage of planet formation, the depletion of the gas disk will increase the eccentricity of  a planet through the sweeping of secular resonance (Nagasawa et al. 2003).  For multiple planetary systems, there are mainly several scenarios that will excite the eccentricities of the planets:  (i) Convergent migration and resonance trap between two planets. A planet in a gaseous disk will migrate inward either due to the imbalance of LR torques  or co-evolute with the viscous disk when it is massive enough to open a gap around it. In the case that the migration speed of inner planet is slower than that of the outer one, or  the inner one is stalled due to the clear of nearby gas, a trap into MMR between the two planets is possible, which may result in the increase of both eccentricities (Lee \\& Peale 2002; Kley 2003).  Due to the trap of MMR,   their eccentricities remain oscillating around  moderate values ($\\sim 0.1-0.3 $) at the end of evolution. The configurations of  GJ 876 b-c in 2:1 MMR and  55 Cnc  b-c in 3:1 MMR are believed to be formed in this way.  (ii) Planetary scattering. During the formation stage of planets, protoplanetary cores may undergo close encounters with the planets, causing ejections of the cores  and the eccentricities excitation for the survival planets. Secular interactions between the survival planets ($m_b$ and $m_c$) may result in oscillations of eccentricities between 0 and a finite value ($\\sim 0.1$). Such a motion is near the separatrix of libration and circulation of difference of perihelion longitude ($\\Delta \\varpi_{bc}$) in the eccentricity plane ($e_b e_c \\cos \\Delta \\varpi_{bc}$, $e_be_c \\sin \\Delta\\varpi_{bc}$),  so it is called a near-separatrix motion (Barnes \\& Greenberg 2006, 2008). This  model can account for  the eccentricity properties of  the $\\upsilon$ Andromedae system (Ford et al. 2005).  (iii) Divergent migration and MMR crossing under interaction with planetesimal disk. After circumstellar disk depletes due to disk accretion, photoevaporation or planet formation within $\\sim 3$ Myrs (Haisch et al. 2001), planets may undergo migration through angular momentum exchanges with residue embryos and planetesimals (Fernandez \\& Ip 1984; Malhotra 1993; Hahn \\& Malhotra 1999). In the solar system, numerical simulations show that, Jupiter will drift inward, while Saturn, Uranus, Neptune may migrate outward, resulting in a divergent migration. During the migration, the cross of 2:1 MMR between Jupiter and Saturn excites the eccentricities of four giant planets (Tsiganis et al. 2005), which may result in the formation of Trojans population of Jupiter and Neptune (Morbidelli et al. 2005), the later heavy bombardment of the terrestrial planets (Gomes et al. 2005), and the architecture of Kuiper belt (Levison et al. 2008).  (iv) Slow diffusion due to planetary secular perturbation. During the final stage of planet evolution when planets are almost formed in well separated orbits, secular perturbations between them result in a slow increase of stochasticity of the system. This procedure can be approximated as a random walk in the space of velocity dispersion, and the  resulting eccentricities of the planets obey a Rayleigh distribution, which agrees with the statistics of  the eccentricities for the observed exoplanets (Zhou et al 2007). The difference between this and the previous planetary scattering scenario is that,  slow diffusion model  may occur in a much longer time span, it is effective especially in the later stage of planet evolution when the planetary orbits are well separated and there is no  violent scattering events.   In this paper, we investigate  the eccentricity formation scenarios and dynamics of the OGLE-06-109L system through N-body simulations, with focuses on the following topics: (a) the origin of the eccentricities for  the two giant planets,  (b)  the stability of the present configuration, (c) the possible existence of planets in  the habitable zone and the outer region. The  eccentricity formation scenarios revealed in this paper  can  be extended to other multiple planetary systems. The paper   is organized as follows. In section 2, the dynamics of the present system is investigated, with much attention paid on the stability of the two giant planets system, the inner and outer regions. Then in section 3, we study the various scenarios (i, ii, iii) as mentioned above to account for the eccentricity formation of the two giant planets. Conclusions and discussions presented in the final section. \\vspace{-1cm} ",
        "conclusions": "In the paper we investigate the dynamics and formation scenario for  OGLE-06-109L system in terms of the observed planets $(m_b,~m_c)$, aiming to understand its formation history and to predict $e_b$ that is not revealed by observation. According to the investigation, we find the evolution history of $m_b$ and $m_c$ depends strongly on the initial conditions, i.e., the cores of $m_b$ and $m_c$ before  efficient gas-accretion begins.   According to the conventional core-accretion scenario of planet formation,  a giant planet forms from a massive embryo ($> 10~ M_\\oplus$) through accreting nearby gas.  Embryos beyond the snow line tend to have larger isolated masses, thus they are the ideal candidates for planetary cores.   However, there will be more than one embryo beyond the snow line. For example,  assuming  a solid disk of 2 times of the minimum mass solar nebula (MMSN)  for the OGLE-06-109L system, the space between 3 AU and  8 AU can be occupied by  4-5   embryos with isolation masses above $10 ~M_\\oplus$ and mutual separations $\\sim 10$ Hill radii. Due to the long quasi-hydrostatic sedimentation stage of gas ($\\sim $ several Myrs, Pollack et al. 1996),  and the perturbation from  first generation giant planets to nearby embryos,  all isolation masses  may have the chance to grow up into second generation giant planets, thus the initial locations of formed planets can not be well determined.  For the two giant planets in  OGLE-06-109L system, if they formed from embryos with relatively far mutual distances, their  initial configuration is  loose, e.g., $a_b/a_c  <  0.48$, so that they are beyond the 3:1 MMR. Subsequent smooth migration under disk tide will  result in a 3:1 MMR, provided suitable gas depletion timescale. We did 8 runs in model 1,  all the simulations result in the  two giant planets trapped in 3:1 MMR, with their eccentricities being excited and osculating around $\\sim 0.1$ at the end of evolutions as in Fig.6c.  % Their eccentricities   will be excited and remain oscillating around $\\sim 0.1$ (Fig.6c). % However,  close encounters  with the residue embryos  may occur during their %migration, which will disturb the previous scenario by  shifting %the semi-major axis ratio,  and thus prevent them into 3:1 MMR in %most cases. Close encounters with embryos may excite $e_c$ up to %$0.1$. Subsequent secular evolution between the two planets will %keep $e_b$ oscillating from around 0 to $\\sim 0.06$, the %near-separatrix motion in the eccentricity plane (Fig.7c).  If  the two planets formed from embryos with a relatively small distance,  they may have a compact configuration initially with  $a_b/a_c  >  0.48$. Then planetary scattering among residue embryos $(m_{e1})$ and planets are most probably the major cause of $e_b$ and $e_c$. Among the 900 simulations of model 2 we did, 185 runs ($20.5\\%$) with close encounter events occurred between $m_{e1}$ and one of the planets, with the average values $e_b \\sim 0.058$ and $e_c \\sim 0.085$ (Fig.10).  Only $3.5\\%$ of the 900 runs lead to  the trap (or in the boundary) of  $m_b$ and $m_c$ into 3:1 MMR.  After the gas disk is almost depleted, divergent migration of $m_b$   and $m_c$ caused by the residue embryos and planetesimals  in outer disk may drive $m_b$ and $m_c$ passing through lower order MMRs. According to our simulations of model 3, the crossing of 2:1 MMR is unlikely in the OGLE-06-109L system, as it will excite eccentricities of $m_b$ and $m_c$ up to $0.2-0.3$, and it is easy to eject $m_c$ out of the system. On the other hand, the crossing of 3:1 MMR is likely, which will excite the eccentricities  up to $e_b\\sim 0.06$ and $e_c\\sim 0.10$. However,  from our simulations, the required migration  depends on the mass and radial location of the planetesimal disk. A solid disk with mass enhancement factor $f_d \\ge 2 $ over the minimum solar nebular may be needed, with their inner edge within $7.8$AU. Considering the similarities of solar system and the OGLE-06-109L system, $f_d \\sim  2$ is still possible for the OGLE-06-109L system with a stellar mass $\\sim 0.5~ M_\\odot$.    In summary, all the three models (i) smooth, convergent migration and the trap of 3:1 MMR; (ii) planetary scattering; (iii) divergent migration and the crossing  of 3:1 MMR, $e_b$ and $e_c$ can be excited. However, the probabilities, the conditions and the final outcomes of these three models are different. Smooth and convergent migration in model (i), if it occurs as predicted by the standard model, could result steadily in the trap of 3:1 MMR between $m_b$ and $m_c$, with $e_b,~e_c\\sim 0.1$ all the time. For model (ii), the probability of planetary scattering occurs is $\\sim 20\\%$,  they result in average  $e_b \\sim 0.06$ and $e_c \\sim 0.09$, and $m_b$ is more likely to undergo  a near-separatrix motion in ($e_be_c \\cos \\Delta \\varpi_{bc},~e_be_c\\sin \\varpi_{bc}$), i.e., $e_b$ passing 0 at a secular timescale ($\\sim 0.08$ Myrs as in Fig.15a). The probability for $m_b,~m_c$ 3:1 MMR crossing in model (iii) depends on the mass and extension of residue solid disk, and  will result in $e_b \\in [0.01-0.07]$ and $e_c \\in [0.03-0.1]$, but the variations in these ranges are in a relative shorter timescale, e.g.,  $\\sim $  0.01 Myrs in Fig.15b.  \\begin{figure}[htp] \\centering  \\includegraphics[scale=0.25,height=2in]{fg15.eps}  %\\vspace{-5cm} \\caption{\\footnotesize{Eccentricity evolutions of a typical run in model (ii) and model (iii). Panel (a): zoomed in evolution of figure 7c from 6 Myrs to 7 Myrs, with  $a_b \\approx 2.40$ AU and $a_c \\approx 5.65$ AU. Panel (b): zoomed in evolution of figure 14d from 2.4 Myrs to 2.6 Myrs, with  $a_b \\approx 2.28$ AU and $a_c \\approx 4.73$ AU. }}  \\end{figure}  Some analytical estimations of related timescale is helpful to reveal the different procedures corresponding to eccentricity evolution in model (ii) and (iii). The  timescale for the secular evolution of two planets is given as $2\\pi/|g_1-g_2|$, where $g_1,~g_2$ are the two eigenfrequencies (Murray \\& Dermott 1999, Zhou \\& Sun 2003). This gives  0.076 Myrs and 0.037 Myrs for orbits in Fig.15a and Fig.15b, respectively. In the circular restricted three-body (CRTB) framework, the timescale of a massless body in the 3:1 MMR of a perturber is estimated as $2\\pi/(ne \\sqrt{3\\mu' \\alpha f_d(\\alpha)})$, where $\\mu'=m'/m_*$ is the mass ratio of the perturber,  $n,e$ is the mean motion and eccentricity of the massless body, $f_d(\\alpha)$ is the function of Laplace coefficients of semi-major axis ratio $\\alpha$ (Murray \\& Dermott 1999). Assuming $m_b$ is the perturber, $e=0.04$ gives the $e$-evolution timescale  of $m_c$ in $m_b$'s 3:1 MMR as 0.01 Myrs, although the CRTB model is not a good model here. So the eccentricity evolution of Fig.15a is due to the secular dynamics, while that in Fig.15b is due to the  3:1 MMR crossing. We also find such a 3:1 MMR timescale is kept for these two orbits up to the end of simulation.  So to understand scenarios of  eccentricity formation for the OGLE-06-109L system, we need more detailed information of their orbits. If $m_b$ and $m_c$ are shown by observation that they are in 3:1 MMR, then model (i) should be the most possible scenario. Based on our simulations in section 2.1, we think  the possibility that $m_b,~m_c$ in 3:1 MMR is small, thus model (i) is unlikely. However, either model (ii) or (iii) can not be decided by the present observations. In both models, the averaged  $e_c\\sim 0.09-0.10$, roughly agrees with the observed value $ 0.11_{-0.04}^{+0.17}$, and  predict the most possible value of $e_b \\approx  0.06$. If $m_b$ and $m_c$ are observed to undergo a near separatrix motion in ($e_be_c \\cos \\Delta \\varpi_{bc},~e_be_c\\sin \\varpi_{bc}$) plane in  a timescale of secular motion, then model (ii) is favored for the origin of eccentricities. And we can infer that, after OGLE-06-109L b and c formed, either smooth migration history during the presence of gas disk is  short, or during their migration history, embryos may suffer close encounter with the planets, exciting their eccentricities. However, if $e_b,~e_c$ oscillates in  a timescale of nearby (crossed) MMR, then most probably  model (iii) accounts for the origin of their eccentricities, and based on this  it is possible to predict the extensions of residue disk mass and location through a more detailed study.   For the stability of the OGLE-06-109L, the two giant planets will be stable provided $e_b^2+e_c^2\\le 0.3^2$. According to the formation scenario, super-Earth planets may be formed inside or outside their orbits.  Numerical simulations show  the region a $\\leq 1.5$ AU (including the habitable zone) or a $\\geq 9.7$ AU is stable. Although the habitable zone contains secular resonance of the system (Migaszewski et al. 2009; Malhotra \\& Minton 2008), our investigation shows that it is wide enough for an Earth-mass planet being formed and stable at least 10 Myrs. In the rare cases when the two giant planets are in the 3:1 MMR, the stable region in  inner orbits  is reduced to $a<1.4$ AU, while that in the outer region is enlarged to $a<7.5$ AU.  When extending the analyses  to other multiple-planet systems, as close encounters between residue embryos are common, we expect planetary scattering and the consequent near-separatrix motion of eccentricities among multiple planetary systems are also common, which agrees with the statistics of the multiple exoplanet systems observed (Barnes \\& Greenberg 2006, 2008). Divergent migration is also possible to sculpt the architecture of the multiple planetary systems. Both  mechanisms can account for the presence of modest eccentricities in the multiple planetary systems without necessarily being trapped in MMRs.   We thank Dr. S. Mao for useful discussions, and the anonymous referee for his constructive suggestions. This work is supported by NSFC (10925313, 10833001, 10778603), National Basic Research Program of China (2007CB814800)."
    },
    "0910/0910.3884_arXiv.txt": {
        "abstract": "We present IRAC/MIPS {\\em Spitzer} observations of intermediate-mass stars in the 5 Myr old {\\LOri} cluster. In a representative sample of stars earlier than F5 (29 stars), we find a population of 9 stars with a varying degree of moderate 24{\\micron} excess comparable to those produced by debris disks in older stellar groups. As expected in debris disks systems, those stars do not exhibit emission lines in their optical spectra. We also include in our study the star HD 245185, a known Herbig Ae object which displays excesses in all {\\em Spitzer} bands and shows emission lines in its spectrum. We compare the disk population in the {\\LOri} cluster with the disk census in other stellar groups studied using similar methods to detect and characterize their disks and spanning a range of ages from 3 Myr to 10 Myr. We find that for stellar groups of 5 Myr or older the observed disk frequency in intermediate mass stars (with spectral types from late B to early F) is higher than in low mass stars (with spectral types K and M). This is in contradiction with the observed trend for primordial disks evolution, in which stars with higher stellar masses dissipate their primordial disks faster. At 3 Myr the observed disk frequency in intermediate mass stars is still lower than for low mass stars indicating that second generation dusty disks start to dominate the disk population at 5 Myr for intermediate mass stars. This result agrees with recent models of evolution of solids in the region of the disk where icy objects form ($>30$ AU), which  suggest that at 5-10 Myr collisions start to produce large amount of dust during the transition from runaway to oligarchic growth (reaching  sizes of $\\sim$500 km) and then dust production peaks at 10-30 Myr, when objects reach their maximum sizes ($\\ge$1000 km). ",
        "introduction": "\\label{sec:int}  Planets are thought to form in circumstellar disks of dust and gas that result from the collapse of rotating molecular cloud cores during the star forming process. These primordial optically thick disks evolve due to viscous processes by which the disk is accreting material onto the star while expanding to conserve angular momentum \\citep{hartmann98}. In addition, dust grains in the disk are expected to grow in size and settle toward the mid-plane of the disk \\citep{weidenschilling97}, where planetesimals and planets can be formed \\citep{backman93,kenyon09}. The dispersal of primordial disks operate less efficiently for stars with lower masses \\citep{muzerolle03, aurora05, lada06, carpenter06, hernandez07a, kennedy09}. Particularly, for low mass stars (K5 or later), 90\\% of the stars lose their primordial disks by about 5-7 Myr \\citep[e.g.][]{haisch01,hernandez08}, while for objects in the mass range of Herbig Ae/Be (HAeBe; types B, A or F early) stars, the corresponding time scale for primordial disk dissipation is less than 3 Myr \\citep{hernandez05}.   The HAeBe stars, which have primordial optically thick disks, are the precursors of intermediate mass stars with debris disks (like Vega or $\\beta$ Pic). Since the estimated lifetime for circumstellar dust grains due to radiation pressure, sublimation, and  Poynting-Robertson drag is significantly smaller than the estimated ages for the stellar systems, the gas-poor disks observed in Vega type objects must be replenished from a reservoir, such as sublimation of comets or collisions between parent bodies \\citep{chen06, kenyon04,dominik03}. Most observational studies of the evolution of debris disks have focused on the long time scales \\citep[e.g][]{decin03, rieke05, su06, siegler07, carpenter09}. These studies show that the luminosities of the debris disks decay following a simple power law (with characteristic time-scales of several hundred millions years), as the debris reservoir of collisional parent bodies diminishes and the dust is removed \\citep{kenyon04,dominik03}. Observations suggest a peak in the debris disk phenomenon in intermediate mass stars at 10-30 Myr \\citep{hernandez06,currie08a}; at the peak, debris disks are more frequent  and have larger infrared excesses. These observational results are supported by theoretical models of evolution of solids in the disk \\citep{kenyon04,kenyon08}. Since debris disks may trace active planet formation \\citep{greenberg78,backman93,kenyon08,wyatt08,cieza08}, studies of debris disks younger than 10-30 Myr are crucial to understanding the processes in which primordial disks evolve to planetary systems.   The {\\LOri} cluster (also known as Collinder 69) represents a unique laboratory for studies of disk evolution because it is reasonably near and relatively populous, making  possible statistically significant studies of disk properties in a wide range of stellar masses. Based on main sequence fitting of early-type stars, \\citet{murdin77} obtained a distance of 400$\\pm$40 pc and age of 4 Myr. Similarly, using Stromgren photometry and main sequence fitting, \\citet{dolan01} derived a distance of 450$\\pm$50 pc and a turnoff age of 6 Myr. Recently, \\citet{mayne08} reported that the {\\LOri} cluster and the $\\sigma$ Orionis cluster have an age of 3 Myr. However, their disk frequencies \\citep{barrado07,hernandez07a} and comparison between empirical isochrones \\citep{hernandez08} suggest that the {\\LOri} cluster is older than the $\\sigma$ Orionis cluster. For consistency with previous studies of disk and stellar population, we assume a distance of 450 pc and an age of 5 Myr \\citep[see][]{mathieu08}.   The unprecedented sensitivity and spatial resolution provided by the {\\em Spitzer Space Telescope} in the near- and mid-infrared windows are powerful tools to expand significantly our understanding of star and planet formation processes. In this contribution, we analyze the near- and mid-infrared  properties of intermediate mass stars in the $\\lambda$ Orionis cluster. This paper is organized as follows. In \\S \\ref{sec:obs} we describe the observational data and the selection of intermediate mass star in the cluster. In \\S \\ref{disk_type} we present the method for identifying stars with infrared excesses. In \\S \\ref{disk_comp} we compare the disk population of the {\\LOri} cluster  with results from other stellar groups inferring the nature of disk emissions around intermediate mass stars. Finally, we give our conclusions in \\S \\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc}  We have used the IRAC and MIPS instruments on board the {\\em Spitzer Space Telescope} to study the frequencies and properties of disks around the intermediate mass population of the 5 Myr old $\\lambda$ Orionis cluster. We have selected 29 intermediate mass stars ($\\sim$F5 or earlier) using optical and 2MASS photometry. One star, HD 245185, displays excesses in all IRAC band, strong excess at 24 {\\micron} and emission line in its optical spectra, supporting previous results in which HD 245185 is a HAe star having an accreting primordial disk \\citep{finkenzeller84, herbst99, winter01, hernandez04, acke05, wade07}. Additionally, we have identified nine stars as debris disks, which exhibit modest excess at 24{\\micron} and very small or no excesses in the IRAC bands. The SEDs of our debris disks sample are comparable to those observed in debris disk populations of older stellar groups. As expected in stars with second generation disks, these stars do not show detectable accretion indicators in their optical spectra.  We have combined observations of the {\\LOri} cluster with those of other stellar groups of ages raging from 3 to 10 Myr located at $<$ 500 pc from the sun, studied using similar methods to detect and characterize their disks. Because of observational thresholds, the observed disk frequencies observed in K and M stars reflect the primordial disk populations, while the observed disk frequencies in intermediate mass stars encompass both primordial and debris disks. As expected in primordial disk evolution the observed frequencies in K and M stars are smaller in older stellar groups. However, observed disk frequencies in intermediate mass stars, represented by a bin in spectral type from B8 to F0, show a behavior in sharp contrast with that expected for primordial disk evolution. In particular, the disk frequency in intermediate mass star increases from $\\sim$20\\% at $\\sim$3 Myr, to $\\sim$40\\% at 5 Myr and to $\\sim$50\\% at $\\sim$10 Myr. Since the timescale for primordial disk dissipation is dependent on stellar mass, the expected frequency of primordial disks around intermediate mass stars at ages $\\ge$ 5 Myr, as inferred from the corresponding frequency in low mass stars, is very low. In stellar population of 5 Myr or older, the observed disk frequencies in intermediate mass stars are larger than in low mass stars. Thus, the increasing in the observed disk frequency at these ages indicates that these disks are dominated by second generation dust. This result agrees with models of evolution of solids in the disks \\citep{kenyon08}, in which at 5-10 Myr collisions start to produce copious amount of dust during the transition from runaway to oligarchic growth. When oligarchs reach sizes $\\sim$500 km, collisions between 1-10 km objects produce debris instead of mergers and the leftover planetesimals are ground to dust."
    },
    "0910/0910.2398_arXiv.txt": {
        "abstract": "We report on new results on the development activity of broad band Laue lenses for hard X-/gamma-ray astronomy (70/100-600 keV). After the development of a first prototype, whose performance was presented at the SPIE conference on Astronomical Telescopes held last year in Marseille (Frontera et al. 2008), we have improved the lens assembling technology. We present the the development status of the new lens prototype that is on the way to be assembled. ",
        "introduction": "\\label{s:intro}  Along with the hard X-ray astronomy up to 70/100 keV, also the gamma--ray astronomy above this energy is moving from direct sky-viewing telescopes to focusing telescopes. With the advent of focusing telescopes in this energy range, a big leap is expected, in both sensitivity and angular resolution. As far as the sensitivity is concerned, the expected increase is by a factor of 100-1000 with respect to the best non-focusing instruments of the current generation (e.g., BeppoSAX/PDS, Ref.~\\citenum{Frontera97}; INTEGRAL/IBIS, Ref.~\\citenum{Ubertini03}). Concerning the angular resolution, the increase is expected by more than a factor 10 (from $\\sim 10$ arcmin of the mask telescopes like INTEGRAL IBIS to less than 1 arcmin). The next generation of  gamma--ray ($>$70/100 keV) focusing telescopes will make use of the Bragg diffraction technique from mosaic-like crystals in transmission configuration (Laue lenses). The expected astrophysical issues that are expected to be solved with the advent of these telescopes are many and of fundamental importance. A discussion of them is done in the context of the mission proposal {\\em Gamma Ray Imager} (GRI), submitted to ESA in response to the first AO of the 'Cosmic Vision 2015--2025' plan (Ref.~\\citenum{Knodlseder07}). For the astrophysical importance of the $>$100 keV band, see also Refs.~\\citenum{Frontera05a,Frontera06,Knodlseder06,Knodlseder07}.  Here we report on the status of our project HAXTEL (= HArd X-ray TELescope) devoted to  developing the technology for building Laue lenses with broad energy passband (70/100--600 keV). The results of this development activity over the last few years have been reported and discussed (see Refs.~\\citenum{Frontera06,Frontera07} and references therein) with a summary of them given in Ref.~\\citenum{Frontera08}.  Last year we reported the first lens Prototype Model (PM) and its performance (see Ref.\\citenum{Frontera08}). Before discussing the activity performed this year, we summarize the main features of the assembly technique and the  performance of the first PM. ",
        "conclusions": "After the development and testing of the first Laue lens prototype (see Ref.~\\citenum{Frontera08}), we have devoted a new effort to improve the lens focusing quality. In this paper we have described the new upgrading of the lens assembling system, that is being used for building a new lens prototype. This is expected to be ready in one--two months. The test results will be reported soon."
    },
    "0910/0910.2351_arXiv.txt": {
        "abstract": "One possible channel for the formation of dwarf galaxies involves birth in the tidal tails of interacting galaxies. We report the detection of a bright UV tidal tail and several young tidal dwarf galaxy (TDG) candidates in the post-merger galaxy NGC\\,4922 in the Coma cluster. Based on a two-component population model (combining young and old stellar populations), we find that the light of tidal tail predominantly comes from young stars (a few Myr old). The {\\it Galaxy Evolution Explorer} ({\\it GALEX}) ultraviolet data played a critical role in the parameter (age and mass) estimation. Our stellar mass estimates of the TDG candidates are $\\sim10^{6-7} M_{\\sun}$, typical for dwarf galaxies. ",
        "introduction": "A tidal dwarf galaxy (hereafter TDG) is defined as ``a dwarf-sized self-gravitating object assembled from tidal debris'' \\citep{hib05}. The formation of dwarf galaxies from tidal interactions was already suggested by \\citet{zwi56}. Two decades later, blue sources were found in the tidal tail of NGC\\, 4038/39 (``the Antennae''), as well as H{\\scriptsize I} tails in several other interacting galaxies \\citep{sch78}. In the 1990s, more TDGs were discovered \\citep{sch90,mir91,mir92}.  During the last two decades, about 20 more interacting galaxies and galaxy groups were investigated for the formation of TDGs \\citep{hib94,yos94,duc97,duc98,duc00,hei00,wei00,hib01,men01,wei03, tem03,mun04,hib05,nef05,boq07,boq09,han09,kor09,smi09,kon09}. TDGs were found to feature blue optical colors revealing the existence of young stellar populations. They are found at the tip of tidal tails, which usually correlates with H{\\scriptsize I} gas density peaks. They also show higher metallicity than typical dwarf spheroidal galaxies, as expected, since they are assembled from recycled materials ejected by their merging parents.  The formation of tidal tails during the encounter of disk galaxies has been reproduced by numerical simulations \\citep[e.g.,][]{too72,bar88,wet07}. \\citet{wal90} explored the triggering of star formation in tidal tails by combining a dynamical model with a code tracking the color evolution. Soon after, more detailed numerical simulations of the formation of dwarf galaxies in tidal tails followed \\citep{bar92,elm93}. More recent simulations achieve much higher resolutions and include chemical evolution \\citep[e.g.,][]{bou08,rec07}. For instance, \\citet{bou08} achieve resolutions high enough to probe down to the masses of star clusters ($\\gtrsim 10^5 M_\\sun$).  Nevertheless, much debate still exists on the fraction of dwarf galaxies born from tidal interactions, because it is difficult to identify TDGs once the tidal tail fades away. We can only identify TDGs by the presence of tidal tails or by a spectroscopic survey around galaxy mergers. Within the standard hierarchical formation scenario, fewer young tidal dwarf galaxies should be expected at present, compared to higher redshifts. Consequently, the number of confirmed TDGs is too small to characterize the population or to provide information for numerical studies.  Here, we focus on the formation of TDGs in the tidal tail of NGC\\,4922.  NGC\\,4922 (M$_{\\rm r} = -22.1$, z$\\sim0.0235$) is a post-merger galaxy located at the outskirts of the Coma cluster. It is known as a merger between an early-type galaxy and a spiral galaxy. Two nuclei from its progenitors can be resolved, separated by $\\sim$10~kpc. This galaxy is especially interesting because the northern nucleus is a weak active galactic nucleus (AGN), probably as a consequence of the merging event. \\citet{and86} and \\citet{alo99} identified the northern nucleus as a Seyfert 2 galaxy with low ionization level. \\citet{alo99} also showed the extended soft X-ray emission which did not originate from the AGN, and possibly relates to ongoing star formation.   In this paper, we present and discuss a UV-bright tidal tail of NGC\\,4922 and TDG candidates obtained from data taken by the  {\\it Galaxy Evolution Explorer} ({\\it GALEX}). In Section~\\ref{sec:obs}, we describe the {\\it GALEX} observations and data reduction with supporting optical imaging and data reduction. In Section~\\ref{sec:model}, we present the synthetic stellar population modeling. Star formation histories on the tidal tail and the TDG candidates are quantified and discussed in Section~\\ref{sec:discussion}. We summarize our findings in Section~\\ref{sec:summary}.  \\begin{figure*}[!ht] \\includegraphics[scale=0.44]{fig1.eps} \\caption{Left: optical $V$ band deep image taken at Lowell Observatory. In addition to the UV-bright tidal tail, we can see faint tidal structures around NGC\\,4922. Right: {\\it GALEX} NUV surface brightness contours are imposed over the optical SDSS r band image. Optical objects and NUV peaks are indicated and contours are given for $\\mu_{\\rm NUV}\\sim$26.0, 27.0, 28.0~mag/arcsec$^2$. The inset shows optical counterparts of `TDG1' and `NUV1' in the SDSS r band. Contours are drawn for $\\mu_{\\rm NUV}\\sim$26.5, 27.0, 28.0~mag/arcsec$^2$ in this case to see the faint object corresponding to `NUV1'. Because of the poor resolution of the {\\it GALEX} data, we could not deblend the NUV light into the two optical objects. \\label{uvknots}} \\end{figure*} ",
        "conclusions": "\\label{sec:discussion}   \\begin{figure}[!ht] \\includegraphics[scale=0.4]{fig5.eps} \\caption{Best fit models from the pixel SEDs of ``NUV1p'' and ``TDG1p'' in the left panels. A dotted line shows the model SED of young component, while the dashed line indicates the contribution of an old component. In the right panels, $\\chi^2$ contours are presented for the 68\\% and 90\\% confidence levels (dashed and solid lines, respectively). A cross indicates the location where a minimum $\\chi^2$ was found in the parameter space. The best-fit values are presented in Table~\\ref{result_tdg}. \\label{pixelgraph}} \\end{figure}   \\begin{figure}[!t] \\includegraphics[scale=0.45]{fig6.eps} \\caption{Result of stellar population modeling with SEDs derived from aperture photometry. Since we could not deblend the UV light around ``TDG1'' and ``NUV1'' into two separate sources, we add the optical SEDs from both objects. The upper panels show results for ``TDG1+NUV,'' the middle panels are for the optical SED of ``TDG1,'' and the bottom panels are for ``NUV1,'' In the left panels dotted lines indicate the contribution of a young component while the dashed line shows a SED of an old component. In the right panels, $\\chi^2$ contours are presented for 68\\% and 90\\% confidence levels (dashed and solid lines, respectively). The result shows young populations of $\\sim5$~Myr. \\label{tdgsed}} \\end{figure}   \\begin{figure}[!t] \\includegraphics[scale=0.4]{fig7.eps} \\caption{Best SSP model fits for ``TDG1+NUV1.'' The best fit is achieved for $t_{\\rm SSP}=3$~Myr with the UV-optical SED, while it is 16 Myr when only the optical data are used to constrain the models (albeit with a large degeneracy). It shows the power of UV data to constrain the age derivation of young stellar populations. \\label{nouv}} \\end{figure}  The multiband data enable us to derive the stellar contents of this system. Note in Figure~\\ref{uvknots} that the UV light along the tail is brightest in the farthest half of the (optical) tail. Furthermore, the brightest UV knot is found at the tip of the tail. It seems that star formation is more active in the outer regions of tidal tails. This finding is consistent with the observational evidence of TDGs at the tip of the tidal tail in the Antennae galaxy \\citep[NGC4038/39,][]{hib05} as well as the theoretical predictions \\citep[e.g.,][]{bou06}.   We found three optical objects which appear to be related with the NUV surface brightness distribution along the tail.  Optical sources around the tidal tail are indicated as ``TDG1,'' ``TDG2,'' and ``TDG3.'' However, note that none of them is located exactly at the peaks in the NUV light.  Hence we label the three NUV peaks separately as ``NUV1,'' ``NUV2,'' and ``NUV3.''  We follow two complementary approaches to determine the color distribution of the sources on the tidal tail. In Section~\\ref{sec:pbp}, we present a pixel-by-pixel analysis, which avoids the ambiguity of matching objects in UV and optical data. Furthermore, we consider aperture photometry of the targets, which increases the signal-to-noise ratio of the photometric data in Section~\\ref{sec:ap}.  \\subsection{Pixel-by-Pixel Analysis of the Tidal Tail Region} \\label{sec:pbp}  Figure~\\ref{colormap} shows the ``FUV $-$ NUV'' and ``NUV $-$ {\\it r}'' color maps derived from pixel magnitude maps in the tidal tail region. Only the pixels with S/N$_{\\rm FUV} > 1$ are shown in ``FUV $-$ NUV'' color map while the pixels with S/N$_{\\rm NUV} > 1$ are presented in ``NUV $-$ {\\it r}'' color map. Surface brightness contours were superimposed on the color maps for $\\mu_{\\rm NUV}=26.0, 27.0, 28.0$~mag/arcsec$^2$. The tidal tail shows blue colors in both cases, suggesting the presence of stellar populations younger than 1~Gyr (NUV $-$ {\\it r} $\\sim$ 4 indicates a SSP of 1~Gyr at solar metallicity).  A photometric SED is constructed for each pixel using the pixel magnitude maps. Stellar population modeling was applied to the pixel SEDs according to the method described in Section~\\ref{sec:model}. The results of the two-component (with $t_{\\rm OC} =12$Gyr) and SSP modeling are presented in Figures~\\ref{pixelmap} and \\ref{pixelmap_ssp}. The result for the ``central'' pixel of each source is also presented in Table~\\ref{result_tdg} (in order to avoid confusion between results using pixel SEDs and SEDs based on aperture photometry, we append a ``p'' at the end of their ID). The ``central'' pixel means a geometric center for TDG 1, 2, 3 and the NUV brightest pixel for NUV 1, 2, 3 actually. $1\\sigma$ errors are presented in Table~\\ref{result_tdg} also. They indicate that the age and internal extinctions are hardly defined in SSP modeling.   The results suggest that star formation has taken place along the tidal tail in the last few hundred Myr. The whole population of ``NUV1p'' is young, about 5 Myr old, with a large dust extinction. All other pixels, except ``TDG3p'', also show significant mass fractions of young stars of 7--200 Myr old. Figure~\\ref{pixelgraph} shows the best model fits for ``NUV1p'' and ``TDG1p''. $\\chi^2$ contours at the 68 and 90 percent confidence level are also presented in the right panels showing that parameters are reasonably well constrained by the data. The open contours of the confidence levels for the ``NUV1p'' source (top right panel) implies that an upper limit to the age of the young component (around 5~Myr) can be inferred. On the other hand, the two-component modeling of ``TDG3p'' shows a very low mass fraction in young stars (0.01 percent), and an old stellar population of a few Gyr for the SSP modeling. Considering that other parts of the tail are much younger, it is possible that ``TDG3'' is not associated with NGC\\,4922.  \\subsection{Aperture Photometry of TDG Candidates} \\label{sec:ap}  \\begin{deluxetable}{cr@{}@{}lr@{}@{}lr@{}@{}lr@{}@{}lr@{}@{}lcc} \\tabletypesize{\\small} \\tablecaption{Pixel-by-pixel modeling results for the central pixels of our 6 UV/optical sources. \\label{result_tdg}} \\tablehead{ \\colhead{ID} &\\multicolumn{2}{c}{$t_{YC,2CM}$} & \\multicolumn{2}{c}{$f_{YC}$} & \\multicolumn{2}{c}{$E(B-V)_{2CM}$} & \\multicolumn{2}{c}{$t_{\\rm SSP}$} & \\multicolumn{2}{c}{$E(B-V)_{\\rm SSP}$} & \\colhead{RA(J2000)} & \\colhead{Dec(J2000)}\\\\ \\colhead{} &\\multicolumn{2}{c}{{\\scriptsize (Myr)}} & \\multicolumn{2}{c}{{\\scriptsize (\\%)}} & \\multicolumn{2}{c}{} & \\multicolumn{2}{c}{{\\scriptsize (Myr)}} & \\multicolumn{2}{c}{} & \\colhead{{\\scriptsize (hh:mm:ss)}} & \\colhead{{\\scriptsize (dd:mm:ss)}} } \\startdata TDG1p & 200 & $^{+110}_{-49.0}$ & 6.5 & $^{+4.3}_{-2.5}$ & 0. & $^{+0.5}$ & 100 & $^{+143}_{-97.4}$ & 0.19 & $^{+0.48}_{-0.064}$ & 13:01:26.62 & 29:17:12.5 \\\\ TDG2p & 7 & $^{+0.02}_{-0.01}$ & 4.5 & $^{+96}_{-2.9}$ & 0.6 & $^{+0.02}_{-0.2}$ & 7 & $^{+1019}_{-0.01}$ & 0.6 & $^{+0.07}_{-0.3}$ & 13:01:26.56 & 29:17:39.4 \\\\ TDG3p & 8 & $^{+420}_{-7}$ & 0.01 & $^{+0.2}$ & 0.05 & $^{+0.07}_{-0.05}$ & $2\\times10^3$ & $^{+1\\times10^4}_{-1\\times10^3}$ & 0.2 & $^{+0.3}_{-0.2}$ & 13:01:25.79 & 29:17:07.5\\\\ NUV1p & 5 & $^{+0.2}_{-4}$ & 100. & $_{-94.4}$ & 0.6 & $^{+0.06}_{-0.5}$ & 5 & $^{+123}_{-4}$ & 0.6 & $^{+0.06}_{-0.5}$ & 13:01:26.47 & 29:17:10.5 \\\\ NUV2p & 19 & $^{+5.4}_{-4.1}$ & 9. & $^{+91}_{-6}$ & 0.45 & $^{+0.14}_{-0.32}$ & 40 & $^{+101}_{-37}$ & 0.35 & $^{+0.35}_{-0.11}$ & 13:01:26.97 & 29:17:21.5 \\\\ NUV3p & 200 & $^{+363}_{-110}$ & 2.5 & $^{+5.5}_{-1.5}$ & 0.005 & $^{+0.6}_{-0.005}$ & 6 & $^{+1051}_{-5}$ & 0.6 & $^{+0.1}_{-0.6}$ & 13:01:26.97 & 29:17:36.5 \\\\ \\enddata  \\end{deluxetable}   \\begin{deluxetable}{cr@{}@{}lr@{}@{}lr@{}@{}lcr@{}@{}lr@{}@{}lc} \\tablecaption{Aperture photometry modeling results for ``TDG1'' and ``NUV1'' \\label{result_tdg2}} \\tablehead{ \\colhead{ID} & \\multicolumn{2}{c}{$t_{YC,2CM}$} & \\multicolumn{2}{c}{$f_{YC}$} & \\multicolumn{2}{c}{$E(B-V)_{2CM}$} & \\colhead{$M_{2CM}$} & \\multicolumn{2}{c}{$t_{\\rm SSP}$} & \\multicolumn{2}{c}{$E(B-V)_{\\rm SSP}$} & \\colhead{$M_{\\rm SSP}$} \\\\ \\colhead{} & \\multicolumn{2}{c}{{\\scriptsize (Myr)}} & \\multicolumn{2}{c}{{\\scriptsize (\\%)}} & \\multicolumn{2}{c}{} & \\colhead{{\\scriptsize ($M_{\\sun}$)}} & \\multicolumn{2}{c}{{\\scriptsize (Myr)}} & \\multicolumn{2}{c}{} & \\colhead{{\\scriptsize ($M_{\\sun}$)}} } \\startdata TDG1+NUV1 & 4 & $^{+1}_{-3}$ & 20. & $^{+80}_{-17}$ & 0.3 & $^{+0.01}_{-0.1}$ & 4.3$\\times10^{6}$ & 3 & $^{+3}_{-2}$ & 0.3 & $^{+0.01}_{-0.1}$ & 9.3$\\times10^{5}$ \\\\ TDG1 & 5 & $^{+0.07}_{-4}$ & 1.2 & $^{+1.8}_{-0.29}$ & 0.16 & $^{+0.024}_{-0.028}$ & 3.1$\\times10^{7}$ & 12 & $^{+1.2}_{-11}$ & 0.09 & $^{+0.1}_{-0.08}$ & 1.1$\\times10^{6}$ \\\\ NUV1 & 5 & $^{+0.06}_{-4}$ & 7. & $^{+93}_{-5}$ & 0.09 & $^{+0.03}_{-0.03}$ & 3.4$\\times10^{6}$ & 5 & $^{+0.06}_{-4}$ & 0.1 & $^{+0.03}_{-0.03}$ & 3.8$\\times10^{5}$ \\\\ \\enddata \\end{deluxetable}    We also carried out aperture photometry of the TDG candidates in the tidal tail in order to achieve a higher signal-to-noise ratio in their photometry. However, ``TDG2'' and ``TDG3'' do not have clear counterparts in the NUV images, they are just embedded in the extended UV structure. Therefore, we cannot apply aperture photometry for those two objects. In the case of ``TDG1'', its UV light is blended with ``NUV1'' and we can see a faint optical counterpart for ``NUV1'' (the inset image in the right panel of Figure~\\ref{uvknots}). Both of them show very blue optical colors ({\\it g - r} = 0.12 and -0.26 for ``TDG1'' and ``NUV1,'' respectively). In fact, the pixel-by-pixel analysis in Section~\\ref{sec:pbp} showed that ``NUV1p'' has significantly younger stellar populations ($\\sim 5$ Myr) than ``TDG1p'' ($\\sim 200$ Myr). ``NUV1'' might be significantly contaminating the UV fluxes of its surrounding regions. Hence, we can not assume that ``TDG1'' is the sole optical counterpart of the brightest UV blob at the tip of the tidal tail.  Since it was impossible to deblend the UV lights between ``TDG1'' and ``NUV1'' at the tip of the tail, we integrated the optical SEDs from both objects in order to match the UV SED. We measured aperture magnitudes of ``TDG1'' and ``NUV1'' with an aperture size, $3 \\times$FWHM in the SDSS data. Also a magnitude of the corresponding UV blob was measured with an aperture of $3 \\times$FWHM from the {\\it GALEX} data. We investigated stellar populations with three kind of SEDs as ``TDG1+NUV1,'' ``TDG1,'' and ``NUV1'' for comparison.  The results of stellar population modeling are presented in Table~\\ref{result_tdg2}. Figure~\\ref{tdgsed} shows the best fits and $\\chi^2$ contours for three cases. The modeling suggests that ``TDG1+NUV1'' has a young stellar population about a few Myr old, contributing a substantial mass fraction up to 100 per cent with considerable internal extinction of $E(B-V)=0.3$~mag. If we consider ``TDG1'' and ``NUV1'' separately, they show the existence of young stellar population of a few Myr old. In all two-component modeling cases the best fits were found when $t_{\\rm OC} = 12$Gyr. The SSP fits suggest similar results, albeit for larger $\\chi^2$ than from two-component fits.  The stellar mass of this system is derived from the $V$-band luminosity estimated from the $r$-band luminosity using a transformation equation extracted from \\citet{jor06}. According to stellar population modeling, the mass of ``TDG1+NUV1'' system was estimated as 4.3 $\\times10^{6} M_{\\sun}$ with two-component modeling and 9.2 $\\times10^{5} M_{\\sun}$ with SSP modeling.  Stellar masses of ``TDG1'' and ``NUV1'' are also calculated separately based on their optical luminosities. Since the stellar mass-to-light ratio increases with the age of the stellar population, the mass of a young population derived by two-component fits is always greater than the mass estimate from SSP fits.  The use of the {\\it GALEX} UV data played a critical role in the spectral fits. For example, Figure~\\ref{nouv} shows that the best SSP models for ``TDG1+NUV1'' indicates the age of 3~Myr, while the fits without UV data lead to 16~Myr. While the $\\chi^2$ distributions for these two estimates are far from Gaussian, they are separated by $3\\sigma$ at least. Recently, \\citet{hib05} and \\citet{nef05} have also demonstrated the ability of {\\it GALEX} data to detect TDGs and to determine the age of young stellar populations (their estimates were based on SSP modeling using a UV color).    Tidal dwarf galaxies represent one of the possible progenitors of the general population of dwarf galaxies, featuring high metallicity and diverse star formation histories. In this paper we study the tidal tail around a post-merger galaxy, NGC\\,4922, in the Coma cluster. Combining UV and optical photometry, we discover a few TDG candidates in the tail. In our analysis, two-component models generally yield a better fit than SSP models for the tidal tail and TDG candidates, although young populations with a single age cannot be ruled out. Also TDG candidates showed a better fit with models with higher metallicity for a young component. According to the pixel-by-pixel and the aperture analyses, two optical TDG candidates and three NUV peaks show predominantly young stellar populations of a few Myr of age, suggesting that recent star formation occurred almost simultaneously along the tidal tail. The result also corresponds with the recent work on the interacting galaxy Arp\\,305 in which the age of TDG candidates was $\\sim$6 Myr \\citep{han09}.  The pixel-by-pixel analysis adopted in this study is relatively new to the field. With the new deep and multiband photometric data, one can now perform statistically meaningful analyses pixel by pixel.  A clear advantage of the pixel-by-pixel analysis is that it allows us to make spatially resolved age maps, instead of estimates for an integrated region. Hence we can investigate the star formation history of the galaxy in connection with the merging event. The Coma cluster is a bit too far for such a detailed investigation. Closer cluster environments, such as Virgo and Fornax, provide better opportunities.  The UV data proved once again to be powerful for searching for young starbursts.  The {\\it GALEX} UV data on hundreds of nearby galaxies, mainly but not exclusively in the NGS survey mode, will provide other interesting opportunities without doubt.    {\\it Facilities:} \\facility{$\\it{GALEX}$}, \\facility{SDSS}, \\facility{Perkins (PRISM)}"
    },
    "0910/0910.2956_arXiv.txt": {
        "abstract": "We have made a search for common proper motion (CPM) companions to the wide binaries in the solar vicinity. We found that the binary GJ 282AB has a very distant CPM companion (NLTT 18149) at a separation $s=1.09 \\arcdeg$. Improved spectral types and radial velocities are obtained, and ages determined for the three components.  The Hipparcos trigonometric parallaxes and the new radial velocities and ages turn out to be very similar for the three stars, and provide strong evidence that they form a physical system. At a projected separation of 55733AU from GJ 282AB, NLTT 18149 ranks among the widest physical companions known. ",
        "introduction": "\\label{sec:intro}  Our long-standing interest in the properties of very wide binaries and their process of dissociation led us to search for additional common proper motions (CPM) companions to the primaries of our catalogue of wide binaries in the solar neighborhood (Poveda  et al. 1994). Other authors have searched for CPM companions to nearby stars (e.g. Lepine \\& Shara 2002, Lepine \\& Bongiorno 2007, Chaname \\& Gould 2004, Makarov et al. 2008, Scholz et al. 2008) but these searches have usually set rather stringent criteria for the acceptable separations and proper motion differences, in order to eliminate, as far as possible, optical companions.  For our purpose we set upper limits to the separations of $1.5\\arcdeg$,  to the difference of the proper motions of $\\vert\\Delta \\mu \\vert \\leq 0.05 \\arcsec$~yr$^{-1}$, and to the difference of the position angles of less than $10\\arcdeg$.  These limits would appear at first sight to be too generous, but we will discuss below some additional criteria to establish the physical nature of the system GJ 282AB - NLTT 18149 in particular.  Here we note only that the extensively studied system composed of Proxima Centauri and $\\alpha$ Centauri AB is believed to be a physical system (see e.g. Wertheimer \\& Laughlin 2006), despite the fact that Proxima has a separation of $2.18\\arcdeg$ from $\\alpha$ Centauri AB, and a difference in the proper motion of $\\vert\\Delta \\mu \\vert = 0.11\\arcsec$~yr$^{-1}$.  Our search produced a handful of interesting systems.  In this paper, we concentrate on GJ 282AB - NLTT 18149, probably the most remarkable one. For the sake of conciseness we shall refer to this system as GJ 282 ABC.  The distant companion, NLTT 18149, has a separation from GJ 282A of $s = 1.09\\arcdeg$ and a difference of proper motion $\\vert \\Delta \\mu \\vert = 0.029\\arcsec$~yr$^{-1}$. We will show that all three components of this system have very similar Hipparcos parallaxes, radial velocities and ages, which, together, constitute strong evidence in favor of GJ 282ABC being a physical system. ",
        "conclusions": "\\label{sec:conc}   The four properties shared by the system GJ 282 AB-NLTT 18149, to wit, common proper motions, common parallaxes, common radial velocities and common ages, constitute  strong arguments in favor of a physical association of the three stars. Nonetheless, the difference in their proper motions $\\vert\\Delta \\mu \\vert = 0.029 \\arcsec$/year is large enough (compared to the error) to be a cause for concern.  There are several possible explanations for this difference. Physically bound systems with large angular separations may have slightly different $\\mu_\\alpha$, $\\mu_\\delta$, $v_r$ simply due to projection effects. Orbital motion of GJ 282AB may also cause such differences. We examine each of these effects in turn. To assess the importance of projection effects we assume that the space motion of G112-29 is identical to that of GJ282A.  Using the standard formulas (e.g. Smart 1962, p.16) we calculate the $\\Delta\\mu_\\alpha$, $\\Delta\\mu_\\delta$, and $\\Delta v_r$ resulting from the different positions and distances of GJ 282A and G112-29. We obtain $\\Delta\\mu_\\alpha = 0.007\\arcsec$~yr$^{-1}$, $\\Delta\\mu_\\delta = 0.004\\arcsec$~yr$^{-1}$ and $\\Delta v_r = 0.23$ km~s$^{-1}$.  These values are much smaller than the observed differences, and hence we conclude that the latter are not due solely to projection effects.  To estimate the importance of orbital motion of the pair GJ 282AB we need the masses of both components. These were calculated from their photometry and the mass-luminosity relation of Reid et al. (2002), and turn ut to be $M_A  = 0.7~M_\\odot$, $M_B = 0.5~M_\\odot$. Assuming the orbit to be circular and perpendicular to the plane of the sky, we obtain a maximum contribution to the radial velocity of $\\Delta v_r = 0.49$ km~s$^{-1}$. The maximum contribution to the tangential velocity is also $0.49$ km~s$^{-1}$, which corresponds to a maximum proper motion of $\\vert\\Delta\\mu\\vert = 0.007\\arcsec$ yr$^{-1}$, much smaller than the observed  $\\vert\\Delta\\mu\\vert = 0.029\\arcsec$ yr$^{-1}$. Conversely, if the orbit is in the plane of the sky, we obtain a maximum contribution to the proper motion of $\\vert\\Delta\\mu\\vert = 0.007\\arcsec$ yr$^{-1}$, again much smaller than the observed $\\Delta\\mu$. In this case, orbital motion contributes nothing to the observed $\\Delta v_r$. We conclude that the orbital motion of GJ 282AB could marginally account for the observed $\\Delta v_r$, but even in the extreme case of an orbit wholly in the plane of the sky cannot account for the observed $\\vert\\Delta\\mu\\vert$. If the orbit is eccentric, we have to multiply these values by at most a factor of 1.4, but the conclusion remains unchanged.  Having shown that both projection effects and orbital motion cannot explain the discrepancy of $\\vert\\Delta\\mu\\vert = 0.029\\arcsec$ yr$^{-1}$ we suggest that the system GJ 282 AB - NLTT 18149 is in the process of dynamical disintegration. It would not be a unique case. There are at least two other multiple systems that appear to be in a similar state. The system $\\theta_{1}$ Orionis ABCD - E (Allen, Poveda \\& Worley 1974; Allen, Poveda \\& Hern\\'andez-Alc\\'antara 2006; S\\'{a}nchez et al. 2008) has been shown to be in the process of ejecting component E, and  the system composed of BN - I - n (Rodr\\'\\i guez et al. 2005, G\\'omez et al. 2006) is also disintegrating.  Thus, GJ 282AB-NLTT 18149 could be another member of the interesting new class of systems caught in the process of gravitational disintegration. A rough calculation shows that component C would have been ejected about 60,000 years ago. This value is much smaller than the estimated ages of the stars, which implies that the dynamical evolution of the hypothetical bound triple ABC proceeded slowly during most of its lifetime, before C finally escaped.  {\\it Acknowledgments.} AP is gratefull to the Direcci\\'{o}n General de Servicios de C\\'{o}mputo Acad\\'{e}mico (DGSCA - UNAM) for the facilities granted. This paper used the Simbad database and IRAF facilities. We are grateful to the anonymous referee for comments that improved this paper."
    },
    "0910/0910.0481_arXiv.txt": {
        "abstract": "{The Berlin Exoplanet Search Telescope II (BEST II) is a small wide field-of-view photometric survey telescope system located at the Observatorio Cerro Armazones, Chile. The high duty cycle combined with excellent observing conditions and millimagnitude photometric precision makes this instrument suitable for ground based support observations for the CoRoT space mission.} {Photometric data of the CoRoT LRa02 target field collected between November $2008$ and March $2009$ were analysed for stellar variability. The presented results will help in the future analysis of the CoRoT data, particularly in additional science programs related to variable stars.} {BEST II observes selected CoRoT target fields ahead of the space mission. The photometric data acquired are searched for stellar variability, periodic variable stars are identified with time series analysis of the obtained stellar light curves.} {We obtained the light curves of $104335$ stars in the CoRoT LRa02 field over $41$ nights. Variability was detected in light curves of $3726$ stars of which \\variables showed a regular period. These stars are, with the exception of $5$ previously known variable stars, new discoveries.} {} ",
        "introduction": "The Berlin Exoplanet Search Telescopes (BEST and BEST II) are two small-aperture wide-field telescope systems operated as ground-based support facilities for the CoRoT space mission (\\cite{baglin}). The first telescope, BEST (\\cite{Rauer2004}, \\cite{Rauer2009}), operational since $2004$, is located at the Observatoire de Haute Provence, France. The second telescope, BEST II, has been operational since April $2007$ at Observatorio Cerro Armazones, Chile. Both telescopes are used to perform preparatory ground-based variability characterization of the target fields of the space mission. Normally the target fields are observed one year ahead of CoRoT.  The CoRoT mission will complete observations of five long-run fields (150 days) and several shorter runs during the nominal mission phase (\\cite{Deleuil2006}). In support of the space mission, variability characterization of three long run and one initial run fields have been performed within the BEST project (\\cite{Karoff2007}, \\cite{Kabath2007}, \\cite{Kabath2008}, \\cite{Kabath2009}), light curves from CoRoT later can be extended with BEST observational data. Thus, significant contributions can be made to the CoRoT additional scientific programs related to variable stars. Usually, several hundred new variable stars of various types ($\\beta$ Lyr, $\\delta$ Sct, Algol type eclipsing binaries, $\\delta$ Cep, RR Lyr and others) are detected easily with the BEST telescopes in the CoRoT fields. In addition, the data from the BEST telescopes can be used during the confirmation process of CoRoT planets, as exemplified by the BEST pre-discovery transit observations of CoRoT 1b and CoRoT 2b (\\cite{Rauer2009}).  The fourth CoRoT long-run field LRa02 was observed with BEST II during the Chilean winter period $2007/2008$, one year prior to the space mission. The acquired light curves were analyzed for stellar variability. \\cite{poretti} have reported new variable stars in the same field. However, their detected variable stars are brighter than the cases presented here. Thus, the BEST II observations of the LRa02 field are complementary to already published results and within the same magnitude range as the CoRoT observations. The detected stars presented here were cross matched with the 2MASS catalog and the light curves of a few known periodic variable stars in the field could be extended by our data.  \\begin{figure} \\centering \\includegraphics[width=8cm, height=8cm]{11909fg1.eps} \\caption{The orientation of BEST II LRa02 subfields with respect to CoRoT's LRa2b field (coordinates J2000.0).} \\label{figfo} \\end{figure} ",
        "conclusions": "We observed CoRoT's LRa02 stellar field with the BEST II telescope during $41$ nights from November $2007$ to March $2008$. The obtained data were searched for periodic variable stars to support CoRoT's additional science programs. In a sample of \\textbf{$104335$} light curves, we identified \\variables stars showing regular periodicity. Almost all of these periodic variable stars are new detections by BEST II with periods ranging from $0.1$ to $35$ days. A classification of the periodic variables has been performed and the members of the resolved classes are specified. In addition, we confirmed variability for $5$ already known stars in the observed field. More detailed information and finding charts will be provided upon request.  \\begin{table} \\caption{Statistics on newly detected variable stars with BEST II in the LRa02 stellar field compared with OGLE.}             % \\label{tabsumm}      % \\centering                          % \\begin{tabular}{l c l }        % \\hline\\hline                 % Property      &  LRa02a & LRa02b \\\\\\hline Nr. of stars detected \t&  37361   \t\t\t& 66974\\\\ Stars with $j>0.5$\t\t& 1858 ($5\\%$)\t\t& 1868 ($3\\%$)\\\\ Periodic variable stars\t& \\varisa ($0.5\\%$)\t& \\varisb ($0.3\\%$)\\\\ DSCT \t\t\t\t&4  \t&3\\\\ DCEP\t\t\t\t&13   &24\\\\ RR Lyr \t\t\t&12 \t&6\\\\ $\\alpha^2$CVn\t&16   &13\\\\ EA(Algol) \t\t&27 \t&20\\\\ EB($\\beta$ Lyr)&16 \t&15\\\\ EW (W UMa)\t\t&19\t&29\\\\ other\t\t\t\t&66 \t&67\\\\ \\hline % \\end{tabular} \\end{table}"
    },
    "0910/0910.2167_arXiv.txt": {
        "abstract": "Inflationary models within string theory exhibit unusual scalar field dynamics involving non-minimal kinetic terms and generically referred to as k-inflation. In this situation, the standard slow-roll approach used to determine the behavior of the primordial cosmological perturbations cannot longer be used. We present a generic method, based on the uniform approximation, to analytically derive the primordial power spectra of scalar and tensor perturbations. At leading order, the scalar spectral index, its running and the tensor-to-scalar ratio are modified by the new dynamics. We provide their new expression, correct previous results at next-to-leading order and clarify the definition of what is the tensor-to-scalar ratio when the sound horizon and Hubble radius are not the same. Finally, we discuss the constraints the parameters encoding the non-minimal kinetic terms have to satisfy, such as the sound speed and the energy scale of k-inflation, in view of the fifth year Wilkinson Microwave Anisotropy Probe (WMAP5) data. ",
        "introduction": "In the context of string theory, cosmic inflation can be achieved through the motion $D$-branes in higher dimensional warped and compact spacetimes~\\cite{Kachru:2003sx}. There, the inflaton appears as the scalar degree of freedom associated with the position of a brane in these extra-dimensions. From a four-dimensional point of view, Lorentz invariance along a $D$-brane necessarily leads to four-dimensional scalar fields exhibiting non-standard kinetic terms, and more precisely of the Dirac--Born--Infeld (DBI) type~\\cite{Dvali:1998pa, Kachru:2003aw, Langlois:2008wt}. In fact, DBI-inflation belongs to the class of k-inflationary models in which the accelerated expansion of the universe can be driven by the scalar field's kinetic terms instead of its potential energy~\\cite{ArmendarizPicon:1999rj}. Assuming the gravity sector to be described by General Relativity, the action of the ''k-inflaton'' $\\varphi(x^\\mu)$ reads \\begin{equation} S=\\dfrac{1}{2\\kappa } \\int \\ud^4x \\sqrt{-g} \\left[ R + 2\\kappa P\\left(X,\\varphi\\right) \\right], \\end{equation} where $\\kappa \\equiv 8\\pi/\\mpl^2$ and $X \\equiv - g^{\\mu \\nu} \\partial_\\mu \\varphi \\partial_\\nu \\varphi /2$. The quantity $P(X,\\varphi)$ is any \\emph{acceptable} functional of $X$ and $\\varphi$ (see Ref.~\\cite{Bruneton:2007si}).  All kinetically modified inflationary models have a ``speed of sound'' \\begin{equation} \\label{eq:csdef} \\cs^2 \\equiv \\dfrac{\\deriv{P}{X}}{\\deriv{P}{X} + 2 X \\dderiv{P}{X}{X}}\\,, \\end{equation} which is generically different from the speed of light. In a flat Friedmann-Lema\\^{\\i}tre--Robertson--Walker (FLRW) universe, it was shown in Ref.~\\cite{Garriga:1999vw}, that the Mukhanov--Sasaki mode function $\\qmode (k,\\eta)= \\zeta \\sqrt{2\\kappa \\epsilon_1}/\\cs$ ($\\zeta$ being the comoving curvature perturbation) satisfies the modified equation of motion \\begin{equation} \\label{eq:qmode} \\qmodeS'' + \\left[\\cs^2(\\eta) k^2 - \\frac{\\nu^2(\\eta)-1/4}{\\eta ^2}\\right] \\qmodeS = 0. \\end{equation} The effective potential is $(\\nu^2-1/4)/\\eta^2 = [\\ln (a \\sqrt{\\epsilon_1}/\\cs)]''$, and all derivatives are with respect to conformal time $\\eta$. The quantity $a(\\eta)$ stands for the FLRW scale factor while $\\epsilon_1 = -\\ud \\ln H/\\ud \\ln a$ is the first Hubble flow function ($H$ being the Hubble parameter). The standard form of this equation is recovered by setting $\\cs=1$ and can be solved by defining a hierarchy of Hubble flow functions encoding the rate of change of the Hubble parameter and its higher order logarithmic derivatives: $\\epsilon_{n+1} \\equiv \\ud \\ln |\\epsilon_n|/\\ud \\ln a$. Assuming slow-roll, i.e. $\\epsilon_n \\ll 1$, one can expand the effective potential $(\\nu^2-1/4)/\\eta^2$ and solve order by order Eq.~(\\ref{eq:qmode}) along the lines of Refs.~\\cite{Schwarz:2001vv, Leach:2002ar, Schwarz:2004tz}. Generalising this method to the k-inflationary case in which $\\cs(\\eta)$ is not constant requires some care. Indeed, both the effective potential and the propagation speed are modified. In the following, we use the uniform approximation to solve Eq.~(\\ref{eq:qmode}) order by order to predict the shape of the tensor and scalar primordial fluctuations at the origin of the CMB anisotropies. ",
        "conclusions": "From the scalar and tensor power spectra given above, one can compare the predicted CMB anisotropies with the current data. In Ref.~\\cite{Lorenz:2008je}, we have performed a Monte--Carlo--Markov--Chains analysis of the WMAP5 data~\\cite{Komatsu:2008hk} against the k-inflationary power spectra. At $95\\%$ of confidence, the flow parameters and the energy scale of k-inflation have to verify \\begin{equation} \\label{eq:nstwosig} 0.003 \\le 2 (\\epsoneP-\\deltaoneP) + (\\epstwoP + \\deltaoneP) \\le 0.075, \\quad \\log (\\epsoneP \\csP) \\le -2.3, \\quad \\ln \\left(10^5 \\dfrac{H_\\piv}{\\mpl} \\right) \\le -0.59 . \\end{equation} Notably, due to the new degree of freedom introduced by $\\cs$, we do not longer find any bound on $\\epsoneP$ alone. The class of k-inflationary models is thus weakly constrained by CMB data. Let us however mention that the subclass of DBI models generate a large amount of non-Gaussianity when $\\cs$ becomes small~\\cite{Chen:2006nt}. In this later case, one can show that the current WMAP5 bounds on non-Gaussianity imposes that $\\log \\epsoneP \\le -1.1$, at two-sigma~\\cite{Lorenz:2008je}.  \\ack It's a pleasure to thank L.~Lorenz and J.~Martin for their comments on the manuscript. This work is supported by the Belgian Federal Office for Science, Technical and Cultural Affairs, under the Inter-university Attraction Pole grant P6/11."
    },
    "0910/0910.2484_arXiv.txt": {
        "abstract": "The Spectrograph for Photometric Imaging with Numeric Reconstruction (SPINR) sounding rocket experiment was launched on 2000 August 4 to record far-ultraviolet (912-1450 \\AA) spectral and spatial information for the giant reflection nebula in the Upper Scorpius region.  The data were divided into three arbitrary bandpasses (\\shortband, \\midband, and \\longband) for which stellar and nebular flux levels were derived.  These flux measurements were used to constrain a radiative transfer model and to determine the dust albedo for the Upper Scorpius region.  The resulting albedos were $0.28\\pm0.07$ for the \\shortband\\ bandpass, $0.33\\pm0.07$ for the \\midband\\ bandpass, and $0.77\\pm0.13$ for the \\longband\\ bandpass. ",
        "introduction": "Dust in the interstellar medium (ISM) is central to many physical processes, such as star and planet formation.  In the ultraviolet (UV), the properties of interstellar dust are largely gleaned from the interactions between dust and the electromagnetic radiation from UV bright stars.  Typically, these interactions manifest themselves through the absorption and scattering of UV light.  The reduction of starlight by dust is commonly referred to as extinction, which has been characterized in the Milky Way Galaxy by an idealized set of wavelength dependent `extinction curves' (see the review by \\citet{fit04}).  The light scattering properties of dust are characterized by two quantities:  the albedo ($a$) and the asymmetry parameter ($g$). Both $a$ and $g$ are functions of the wavelength of the incident light.  Previous studies have probed reflection nebulae, dark clouds, and diffuse galactic light in an attempt to characterize the scattering properties of interstellar dust from infrared (IR) to UV wavelengths (see the review by \\citet{gor04}).  Although these studies have contributed to the characterization of dust in the ISM, there is still much debate about dust radiative transfer models and the appropriateness of techniques used in measuring $a$ and $g$ for a given sight line \\citep{mat02, dra03, zub04}.  Further observations of dusty environments throughout our galaxy and beyond, along with the continued refinement of current dust models are needed in order to enhance our understanding.  The Upper Scorpius cloud is ideally suited for probing the scattering properties of dust grains.  It is an internally illuminated reflection nebula, which makes the total nebular flux strongly dependent on $a$ and only weakly dependent on $g$ \\citep{gor94}.  This nebular geometry will result in an accurate measurement of $a$, but it is likely that any measure of $g$ will have large uncertainties.  The Upper Scorpius region was also found by \\citet{gau93} to have very little foreground dust, which would add to the uncertainty in the scattering properties of the nebula. By considering this region as a whole, a large sample of UV bright stars are available as illumination sources for the dust model.  Multiple illuminating sources will reduce the model's dependence on the accuracy of any one star's derived properties.  It was for these reasons that the nebula in Upper Scorpius was selected as the target for the Spectrograph for Photometric Imaging with Numeric Reconstruction (SPINR) sounding rocket mission (Cook et al. 2003, hereafter Paper I).  The SPINR experiment provides a unique look at the ISM in the FUV wavelength regime from 912-1450 \\AA.  By recording wide-field ($\\sim$16\\degr) spectral and spatial information from a given target simultaneously, a complete picture of the dust properties in a region can be derived.  Paper I presents the details of the SPINR experiment while its data have been used by Lewis et al. (2005, hereafter Paper II) to derive the extinction properties along several lines of sight in the Upper Scorpius Cloud. Paper II found that the Upper Scorpius region possesses a wide range of extinction properties, which makes it an excellent probe of the various dust environments seen throughout the Milky Way.  Thus far, only a handful of studies have probed dust scattering properties shortward of Ly$\\alpha$ (1216 \\AA) \\citep{wit93, bur02, suj05, sha06, suj07}.  The SPINR data are unique in that spectra for the entire 16\\degr\\ field-of-view (FOV) were recorded simultaneously without the FOV limitations of a traditional long slit or the spectral resolution limitations of traditional wide band filters.  Currently, there are significant discrepancies between model predictions of dust albedo in the FUV and observations \\citep{gor04}. The insights gained from the SPINR study of the Upper Scorpius cloud will help to paint a more complete picture of dust scattering properties in the FUV.  The goal of this paper is to present the determination of $a$ and $g$ values for dust in the Upper Scorpius region at the effective wavelengths $\\lambda_{eff}=$ 973 \\AA, 1106 \\AA, and 1261 \\AA.  Additionally, the Upper Scorpius data presented in \\citet{gor94} were revisited with an improved radiative transfer model to determine the value of $a$ in the STS-39 FUV camera (FUVCam) bandpasses, $\\lambda_{eff}=$ 1362 \\AA\\ and 1769 \\AA. ",
        "conclusions": "\\label{dis}  Dust model simulations were run for each of the SPINR wavelength bands to identify the value for the dust albedo, $a$, that produced the best fit to the data.  This was accomplished by finding the $a$ value that reproduced the $F_N/F_*$ for each bandpass while varying the asymmetry parameter, $g$, from $0.0$ to $0.99$.  As shown in Figure \\ref{ag}, there was very little variation in the best-fit value for the albedo over the entire range of $g$ values for all three of the SPINR bandpasses and the two STS-39 FUVCam bandpasses.  The best fit values for the dust albedo were $a=0.28\\pm0.07$ in the \\shortband\\ bandpass, $a=0.33\\pm0.07$ in the \\midband\\ bandpass, and $a=0.77\\pm0.13$ in the \\longband\\ bandpass. Additionally, the STS-39 data presented in \\citet{gor94} were used in the updated radiative transfer model to make new estimates of the dust albedo in the Upper Scorpius region for the STS-39 FUVCam 1230-1600 \\AA 1650-2000 \\AA bandpasses.  The new albedo values that provided the best fit to the STS-39 FUVCam data were slightly modified from the original estimate of $a=0.47-0.70$ for $\\lambda_{eff}=1362 \\AA$ and $a=0.55-0.72$ for $\\lambda_{eff}=1769 \\AA$, to $a=0.54\\pm0.06$ and $a=0.59\\pm0.07$ respectively.  The optical depth of each star in the model Scorpius cloud is determined mostly from the derived extinction curves, but also has some dependence on the values selected for $R_V$ and $\\tau_{\\rm H_2}$ (see section \\ref{tau}). In order to test the strength of the dependence of the derived albedo values on $R_V$ and $\\tau_{\\rm H_2}$, simulations were run for models with no H$_2$ absorption and $R_V=3.1$ for every stellar line of sight. Not accounting for H$_2$ absorption and using a uniform value of $R_V=3.1$ reduced the albedo values by $<$5\\%, which is still well within the quoted uncertainties for $a$.  Dust model simulations were also run where presence of H$_2$ fluorescence in the Upper Scorpius cloud was ignored (see section \\ref{h2em}). Not accounting for H$_2$ emission in the nebular data results in a $\\sim$40\\% increase in $F_N/F_*$.  This increase in $F_N/F_*$ resulted in values for the albedo of $a=0.35\\pm0.05$ in the \\shortband\\ bandpass, $a=0.43\\pm0.05$ in the \\midband\\ bandpass, and $a=0.93\\pm0.04$ in the \\longband\\ bandpass, which are all within 2$\\sigma$ of the best albedo estimates accounting for H$_2$ emission.  Overall, the derived albedo values are not particularly sensitive to the amount of H$_2$ absorption or $R_V$ value selected along any given line of sight.  However, moderate inconsistencies can be introduced into the albedo values if H$_2$ fluorescence is not considered.  Because SPINR was able to record spatial information about the dust in the Upper Scorpius region in addition to spectral information, it was possible to determine the value of $g$ that best reproduced the SPINR data.  Model images were produced over a range of $g$ values from 0.00 to 0.99 and with $a$ values corresponding to the best fit for each $g$ value (Figure \\ref{ag}). These model images were ``spun'' into sinogram space using a transformation matrix and then compared to the SPINR sinograms using a cross-correlation metric.  The error in the cross-correlation metric was determined using a Monte Carlo simulation in which up to $\\pm 1\\sigma$ worth of noise was added to each pixel in the SPINR sinogram.  The optimal $g$ value was determined from the image/sinogram with the highest cross-correlation to the SPINR data.  The range of acceptable $g$ values was determined from the Monte Carlo error estimate on the highest cross-correlation.  All $g$ cross-correlation metrics that fell within the error of the optimal fit were deemed acceptable.  The $g$ value that provided the best fit to the SPINR \\midband\\ bandpass data was $g=0.96^{+0.02}_{-0.18}$.  Figure \\ref{sino} shows sinogram and image data for the SPINR \\midband\\ bandpass, the corresponding best fit dust model($g=0.96$ and $a=0.30$), and a model outside of the range of acceptable $a$ and $g$ values ($g=0.0$ and $a=0.75$).  In the \\shortband\\ and \\longband\\ bandpasses no statistically significant constraints could be put on the $g$ value.  It is not surprising that values for $g$ are not well determined in this study since the geometry of the nebula in Upper Scorpius was selected for its weak dependence on $g$ to better determine $a$.  However, given the constraint from the \\midband\\ bandpass we can confidently say that $g>0.42$ at the three sigma level.   This is not a tight constraint on $g$, but does align with the general expectation for $g$ to be highly forward scattering, $g>0.7$, as shown by previous predictions for $g$ in the FUV \\citep{gor04}.  The results for $a$ from this study are compared with previous studies of reflection nebulae in Figure \\ref{alit}. The values of $a$ derived for the \\shortband\\ and \\midband\\ bandpasses correspond well with both previous observational determinations $a$ and the \\citet{wei01} model in this spectral region.  The value of $a$ determined for the \\longband\\ bandpass is on the high end of the observational data points scattered within the 1200-1400 \\AA\\ range, but are well within $3\\sigma$ of the other observational determination of $a$.  However, what is clear from this study and previous observational studies of albedo in the FUV is that the dust models \\citep{wei01,cla03,zub04} do not correctly account for the apparent rise in the albedo from around 1200 \\AA\\ to 1800 \\AA.  Further observational studies of the ISM in this spectral region are needed to fully characterize this rise in the albedo and determine its source."
    },
    "0910/0910.5571_arXiv.txt": {
        "abstract": "We present radial velocities of the double-lined spectroscopic binary NP\\,Aqr. The radial velocities and the optical light curves obtained by Hipparcos and ASAS-3 are analyzed separately. The masses of the primary and secondary components have been found to be 1.65$\\pm$0.09 and 0.99$\\pm$0.05 M$_{\\odot}$, respectively. The cross-correlation functions indicate triple peaks which show presence of a tertiary star. The spectroscopic properties of this additional component resemble to that of the primary star. The analysis of the light curves yielded that the more massive primary star fills its corresponding Roche lobe. The secondary component is at or near Roche lobe indicating a new $\\beta$ Lyrae-type near-contact binary. The orbital inclination is about 40$^{\\circ}$ and, therefore, the observed light variations are produced only by the proximity effects. Due to the absence of eclipses, the effective temperature of the secondary star and the radii of the components could not be determined accurately. We conclude that NP\\,Aqr is a non-eclipsing double-lined spectroscopic binary with a distance of about 134$\\pm$7 pc. The absolute parameters of the components are also compared with the evolutionary models. While the location of the primary star seems to be suitable with respect to its mass in the Hertzsprung-Russell diagram, the secondary component is located as if a star having a mass less than 0.6 M$_{\\odot}$. This discrepancy is originated from very low effective temperature determined only from the light curve produced by proximity effects. The distance to the third star appears to be very close to that of the close binary which indicates that it may be dynamically bounded to the binary. ",
        "introduction": "The light variability of NP\\,Aquarii (=HD\\,198528=HIP\\,102935=NSV\\,25363) was first suspected by Tobin, Viton \\& Sivan (1994) during the Spacelab-1 Very Wide Field survey of UV-excess objects. This survey has revealed a variety of stellar objects with strong ultraviolet excess (see, for example, Viton et al. 1988). Perryman et al. (1997) confirmed its light variation using the H$_p$ magnitudes obtained by the Hipparcos mission. The type of the light variability could not be classified with the available data. Later on, Otero (2003) gave the epoch of minimum light, the orbital period and the variability class using the ASAS-3 (Pojmanski, 2002) and Hipparcos databases. He classified the star as a $\\beta$ Lyrae-type eclipsing binary with an orbital period of 0.806982 day and with very shallow eclipses, namely 0.1 mag. Very recently, Kazarovets et al. (2006) designated the star HD\\,198528 as NP\\,Aqr using the criteria of {\\it General Catalog of Variable Stars} in the 78$^{th}$ Name-list of Variable Stars. van Leeuwen (2007) re-examined the Hipparcos data and concluded that NP\\,Aqr is an eclipsing binary with a spectral type of F0V at a distance of 186$\\pm$27 pc. He also gives wide-band B-V and V-I color indices and interstellar reddening, E(B-V). Due to its very limited light variations and a spectral type of F0 Handler (1999) included the star in the $\\gamma$ Doradus candidates but with doubts that it may also be classified as a $\\delta$ Scuti type variable. Handler and Shobbrook (2002) made Johnson  B, V and Cousins I$_c$  observations of the known and candidate $\\gamma$ Dor stars to assess the relationship between the $\\gamma$ Dor and $\\delta$ Sct stars. They found slow variations with amplitudes of several hudredths of a magnitude. Taking into account a double-wave light curve they have excluded W UMa-type variability and suggested a possibility of an ellipsoidal variable with a period of about 0.807 day.  In this paper, we use the optical spectra of NP\\,Aqr to reveal the nature of its light variability and physical properties combining with the photometric data obtained by Hipparcos and ASAS-3. The paper is organized as follows. In \\S 2 the spectroscopic observations and data analysis are described. The spectroscopic and photometric results are combined and the absolute parameters of the stars are derived and the structure of NP\\,Aqr is discussed in \\S 3.  A brief conclusion is given in \\S 4. ",
        "conclusions": "The first spectoscopic observations of the relatively short-period close binary NP\\,Aqr are presented. CCFs indicate that NP\\,Aqr is a triple system. The radial velocities of both components are obtained and analyzed for the spectroscopic elements. NP\\,Aqr has been observed also by the Hipparcos mission and ASAS-3 which yielded almost complete light curves of the system. Both light curves are analyzed separately and the resultant orbital elements are determined. The analysis of the light curves revealed that the light variation with an amount of about 0.1 mag is produced from the proximity effects. In other words, NP\\,Aqr is a double-lined, non-eclipsing spectroscopic binary. The radii of the components and effective temperature of the secondary star could be determined with great uncertainty due to the absence of the eclipses and also presence of a third star. Comparing stellar evolutionary models for the mass of the components we find that location of the primary star in the HR diagram is suitable with its mass. However, the secondary star looks like a low-mass star. While the radius of this star is as large as 1.67 times with respect to its mass, the effective temperature is unexpectedly low. The effective temperatures of the binary star's components can be determined by comparing the depths of the eclipses. Since the stars of NP\\,Aqr do not show eclipses due to the small orbital inclination, the eclipses do not set in. Our analysis indicates that the star NP\\,Aqr shows most of the characteristics of the so-called near-contact binaries. Moreover, we detected a third component in the system NP\\,Aqr. The third star looks like the primary component in the spectra. Using this property and assuming that it is a main-sequence star we determined its distance to be 122 pc, at nearly the same distance with the spectroscopic binary."
    },
    "0910/0910.0738_arXiv.txt": {
        "abstract": "Current status and results of the experiment on recording neutrino bursts are presented. The observation time (since 1980) is 24.7 years. The upper bound of collapse frequency in our Galaxy is 0.093 $y^{-1}$ (90\\% CL). \\vspace{1pc} ",
        "introduction": "One of the current task of the Baksan Underground Scintillation Telescope (BUST) is the search for neutrino bursts from gravitational collapse of stars.  BUST is located in the Northern Caucasus in the underground laboratory at the effective depth of $8.5\\times 10^4\\ g\\cdot cm^{-2}$ (850 m of w.e.) \\cite {Alex}. The facility has dimensions $17\\times 17\\times 11$ m$^3$ and consists of four horizontal scintillation planes and four vertical ones (Fig.~1). Five planes of them are external planes and three lower horizontal planes are internal ones. The upper horizontal plane consists of 576 ($24\\times 24$) liquid scintillator detectors of the standard type, three lower planes have 400 ($20\\times 20$) detectors each. The vertical planes have $15\\times 24$ and $15\\times 22$ detectors. The detector sizes are $0.7\\times 0.7\\times 0.3\\ m^3$. The distance between neighboring horizontal scintillation layers is 3.6 m. The angular resolution of the facility is 2$^o$, time resolution is 5 ns.  The information from each detector is transmitted over three channels:  an anode channel (which serves for trigger formation and amplitude measurements up to 2.5 GeV), a pulse channel with operation threshold 12 MeV (since 1991 this threshold = 8 MeV; the most probable energy deposition of a muon in a detector is 50 MeV $\\equiv $ 1~relativistic particle) and a logarithmic channel with a threshold $s_o = 0.5$ GeV.  The signal from the fifth dynode of PM tube FEU-49 goes to a logarithmic channel (LC) where it is converted  into a pulse whose length t is proportional to the logarithm of the amplitude of the signal \\cite {Bak}. \\begin{figure}[!t] \\centering \\includegraphics[width=3in]{icrc0366_fig01} \\caption{General view of the Baksan underground scintillation telescope.} \\label{fig:fig1} \\end{figure}  There is a voluminous set of monitoring programs which allows us to examine the operation of each detector of the facility in any time interval and discard not properly working detectors. ",
        "conclusions": ""
    },
    "0910/0910.1813_arXiv.txt": {
        "abstract": "{} {We have investigated the differences in apparent opening angles between the parsec-scale jets of the active galactic nuclei (AGN) detected by the $Fermi$ Large Area Telescope (LAT) during its first three months of operations and those of non-LAT-detected AGN.} {We used 15.4 GHz VLBA observations of sources from the 2 cm VLBA MOJAVE program, a subset of which comprise the statistically complete flux density limited MOJAVE sample. We determined the apparent opening angles by analyzing transverse jet profiles from the data in the image plane and by applying a model fitting technique to the data in the $(u,v)$ plane. Both methods provided comparable opening angle estimates.} {The apparent opening angles of $\\gamma$-ray bright blazars are preferentially larger than those of $\\gamma$-ray weak sources. At the same time, we have found the two groups to have similar intrinsic opening angle distributions, based on a smaller subset of sources. This suggests that the jets in $\\gamma$-ray bright AGN are oriented at preferentially smaller angles to the line of sight resulting in a stronger relativistic beaming. The intrinsic jet opening angle and bulk flow Lorentz factor are found to be inversely proportional, as predicted by standard models of compact relativistic jets. If a gas dynamical jet acceleration model is assumed, the ratio of the initial pressure of the plasma in the core region $P_0$ to the external pressure $P_\\mathrm{ext}$ lies within the range 1.1 to 34.6, with a best fit estimate of $P_0/P_\\mathrm{ext}\\approx2$.} {} ",
        "introduction": "The EGRET telescope onboard the {\\it Compton Gamma Ray Observatory} \\citep{Hartman99} provided detections of $\\gamma$-ray emission from many extragalactic point sources, most of which were identified with blazars \\citep[e.g.,][]{M01,SRM03}. The latter term generally refers to objects classified as flat-spectrum radio-loud quasars and BL Lacs. Although blazars comprise only a few per cent of the overall AGN population, they dominate the extragalactic high-energy sky. A systematic comparison of parsec-scale radio jet structure with $\\gamma$-ray emission in AGN has hitherto been problematic and inconclusive due to the limited sensitivity and angular resolution of the EGRET instrument. For example, a recent analysis of a large 6~cm VLBA Imaging and Polarimetry Survey \\citep[VIPS;][]{Taylor07} has hinted that EGRET-detected blazars might have larger than average jet opening angles, but very poor statistics (only four EGRET/VIPS sources) did not allow the authors to draw a firm conclusion.  On 2008 June 11, the {\\it Fermi Gamma-Ray Space Telescope} (previously known as {\\it GLAST}) was successfully launched by NASA with the Large Area Telescope (LAT), a successor to EGRET, onboard. The LAT began operating in August 2008 and has provided $\\gamma$-ray observations with greatly improved sensitivity, superior angular resolution, large field of view, and wide energy range from about 20~MeV to over 300~GeV \\citep{LAT}. During the first three months of science operations, the LAT has made 205 high-confidence ($>10\\sigma$) detections of bright $\\gamma$-ray sources, 116 of which were associated with high galactic latitude ($|b|>10\\degr$) AGN (the LAT Bright AGN Sample together with 10 lower-confidence associations: LBAS-ext), as discussed by \\cite{LBAS}. Their VLBI radio source identifications were reported by \\cite{VLBI_LBAS}. The LAT detections of flaring AGN have also triggered a number of multiwavelength (from radio to $\\gamma$-ray) campaigns \\citep[e.g.,][]{3C84_Fermi}.  The {\\it Fermi} era has already heralded a number of important results that link the $\\gamma$-ray emission and radio properties of AGN. It has been shown that LAT-detected blazars are brighter and more luminous in the radio domain at parsec scales \\citep{MF2}, have higher apparent jet speeds \\citep{MF1}. The LAT-detected blazars also seem to have higher Doppler factors determined from the mm-wavelength variability \\citep{MF3}. The $\\gamma$-ray and radio flares are found to appear within a typical timescale of up to a few months and are likely associated with the parsec-scale radio core \\citep{MF2}. In this Letter we continue investigating the relations between $\\gamma$-ray brightness and the properties of parsec-scale radio jets from the LBAS-ext list as probed by interferometric observations with the VLBA. Hereafter we refer to the Fermi LAT 3-month bright source list \\citep{LBAS} when using the term LAT-detected.  Throughout this letter, we use the term ``core'' as the apparent origin of AGN jets that commonly appears as the brightest feature in VLBI images of blazars \\citep[e.g.,][]{L98,Marscher08}. We use the $\\Lambda$CDM cosmological model with $H_0=71$~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_m=0.27$, and $\\Omega_\\Lambda=0.73$. ",
        "conclusions": "\\label{s:summary}  We have measured the projected jet opening angles on parsec scales for 115 blazars (29 LAT-detected and 86 non-LAT-detected) from the complete flux-density limited MOJAVE-1 sample and for 27 additional LAT-detected sources monitored by the MOJAVE program. The apparent opening angles for $\\gamma$-ray bright sources are on average larger than those in $\\gamma$-ray weak ones, while the intrinsic opening angle distributions based on smaller samples are statistically indistinguishable. We interpret this as an evidence for $\\gamma$-ray bright blazars to have preferentially smaller viewing angles and, consequently, stronger relativistic beaming, which boosts emission in both, the $\\gamma$-ray and radio bands. This conclusion is consistent with recently obtained results that show LAT-detected blazars to be brighter and more luminous in radio domain \\citep{MF2}, to have faster jets \\citep{MF1} and higher variability Doppler factors \\citep{MF3}. There is an indication for BL~Lac objects to have on-average wider intrinsic opening angles than those of quasars.   The intrinsic opening angle and Lorentz factor are found to be inversely proportional for a sample of 56 AGN jets in agreement with theoretical predictions of simple gas dynamical and magnetic acceleration models. The best approximation of the inferred relation is achieved assuming a ratio of internal and external jet pressure $P_0/P_\\mathrm{ext}\\approx2$ in gas dynamical models of compact relativistic jets."
    },
    "0910/0910.3211_arXiv.txt": {
        "abstract": "We present six simulations of Galactic stellar haloes formed by the tidal disruption of accreted dwarf galaxies in a fully cosmological setting. Our model is based on the Aquarius project, a suite of high resolution N-body simulations of individual dark matter haloes. We tag subsets of particles in these simulations with stellar populations predicted by the \\galform{} semi-analytic model. Our method self-consistently tracks the dynamical evolution and disruption of satellites from high redshift. The luminosity function and structural properties of surviving satellites, which agree well with observations, suggest that this technique is appropriate. We find that accreted stellar haloes are assembled between $1<z<7$ from less than 5 significant progenitors. These progenitors are old, metal-rich satellites with stellar masses similar to the brightest Milky Way dwarf spheroidals ($10^{7}-10^{8}\\;\\rm{M_{\\sun}}$). In contrast to previous stellar halo simulations, we find that several of these major contributors survive as self-bound systems to the present day. Both the number of these significant progenitors and their infall times are inherently stochastic. This results in great diversity among our stellar haloes, which amplifies small differences between the formation histories of their dark halo hosts. The masses ($\\sim10^{8}-10^{9}\\;\\rm{M_{\\sun}}$) and density/surface-brightness profiles of the stellar haloes (from 10--100~kpc) are consistent with expectations from the Milky Way and M31. Each halo has a complex structure, consisting of well-mixed components, tidal streams, shells and other subcomponents. This structure is not adequately described by smooth models. The central regions ($<10$~kpc) of our haloes are highly prolate ($c/a\\sim0.3$), although we find one example of a massive accreted thick disc. Metallicity gradients in our haloes are typically significant only where the halo is built from a small number of satellites. We contrast the ages and metallicities of halo stars with surviving satellites, finding broad agreement with recent observations. ",
        "introduction": "\\label{sec:introduction}   An extended and diffuse stellar halo envelops the Milky Way. Although only an extremely small fraction of the stars in the Solar neighbourhood belong to this halo, they can be easily recognized by their extreme kinematics and metallicities. Stellar populations with these properties can now be followed to distances in excess of 100~kpc using luminous tracers such as RR Lyraes, blue horizontal branch stars, metal-poor giants and globular clusters (e.g. Oort 1926; Baade 1944; Eggen, Lynden-Bell \\& Sandage 1962; Searle \\& Zinn 1978; Yanny {et~al.} 2000; Vivas \\& Zinn 2006; Morrison {et~al.} 2009).  In recent years, large samples of halo-star velocities (e.g. Morrison {et~al.} 2000; Starkenburg {et~al.} 2009) \\newt{and} `tomographic' photometric and spectroscopic surveys have shown that the stellar halo is not a single smoothly-distributed entity, but instead a superposition of many components (Belokurov {et~al.} 2006; Juri{\\'c} {et~al.} 2008; Bell {et~al.} 2008; Carollo {et~al.} 2007, 2009; Yanny {et~al.} 2009). Notable substructures in the Milky Way halo include the broad stream of stars from the disrupting Sagittarius dwarf galaxy (Ibata, Gilmore \\& Irwin 1994; Ibata {et~al.} 2001), extensive and diffuse overdensities (Juri{\\'c} {et~al.} 2008; Belokurov {et~al.} 2007a; Watkins {et~al.} 2009), the Monoceros `ring' (Newberg {et~al.} 2002; Ibata {et~al.} 2003; Yanny {et~al.} 2003), the orphan stream (Belokurov {et~al.} 2007b) and other kinematically cold debris (Schlaufman {et~al.} 2009). Many of these features remain unclear. At least two kinematically distinct `smooth' halo components have been identified from the motions of stars in the Solar neighbourhood, in addition to one or more `thick disc' components (Carollo {et~al.} 2009). Although current observations only hint at the gross properties of the halo and its substructures, some general properties are well-established: the halo is extensive ($>100\\:\\rm{kpc}$), metal-poor ($\\mathrm{\\langle[Fe/H]\\rangle\\sim{-1.6}}$, e.g. Laird {et~al.} 1988; Carollo {et~al.} 2009) and contains of the order of 0.1-1\\% of the total stellar mass of the Milky Way \\newt{(recent reviews include Freeman \\& Bland-Hawthorn 2002; Helmi 2008)}.  Low surface-brightness features seen in projection around other galaxies aid in the interpretation of the Milky Way's stellar halo, and vice versa. Diffuse concentric `shells' of stars on $100$~kpc scales around otherwise regular elliptical galaxies \\newt{have been} attributed to accretion events (e.g. Schweizer 1980; Quinn 1984). Recent surveys of M31 (e.g. Ferguson {et~al.} 2002; Kalirai {et~al.} 2006; Ibata {et~al.} 2007; McConnachie {et~al.} 2009) have revealed an extensive halo (to $\\sim150\\,\\rm{kpc}$) also displaying abundant substructure. The surroundings of other nearby Milky Way analogues are now being targeted by observations using resolved star counts to reach very low effective surface brightness limits, although as yet no systematic survey has been carried out to sufficient depth. (e.g. Zibetti \\& Ferguson 2004; McConnachie {et~al.} 2006; de~Jong, Radburn-Smith \\&  Sick 2008; Barker {et~al.} 2009; Ibata, Mouhcine, \\& Rejkuba 2009). A handful of deep observations beyond the Local Group suggest that stellar haloes are ubiquitous and diverse (e.g. Sackett {et~al.} 1994; Shang {et~al.} 1998; Malin \\& Hadley 1999; Mart{\\'{i}}nez-Delgado {et~al.} 2008, 2009; Fa{\\'{u}}ndez-Abans {et~al.} 2009).  Stellar haloes formed from the debris of disrupted satellites are a natural byproduct of hierarchical galaxy formation in the \\lcdm{} cosmology\\footnote{In addition to forming components of the accreted stellar halo, infalling satellites may cause dynamical heating of a thin disc formed `in situ' (e.g. Toth \\& Ostriker 1992; Velazquez \\& White 1999; Benson {et~al.} 2004; Kazantzidis {et~al.} 2008) and may also contribute material to an accreted thick disc (Abadi, Navarro \\& Steinmetz 2006) or central bulge. We discuss these additional contributions to the halo, some of which are not included in our modelling, in \\mnsec{sec:defining_haloes}}. The entire assembly history of a galaxy may be encoded in the kinematics, metallicities, ages and spatial distributions of its halo stars. Even though these stars constitute a very small fraction of the total stellar mass, the prospects are good for recovering such information from the haloes of the Milky Way, M31 and even galaxies beyond the Local Group (e.g. Johnston, Hernquist \\& Bolte 1996; Helmi \\& White 1999). In this context, theoretical models can provide useful `blueprints' for interpreting the great diversity of stellar haloes and their various sub-components, and for relating these components to fundamental properties of galaxy formation models. Alongside idealised models of tidal disruption, \\textit{ab initio} stellar halo simulations in realistic cosmological settings are essential for \\newt{direct} comparison with observational data.  In principle, hydrodynamical simulations are well-suited to this task, as they incorporate the dynamics of a baryonic component self-consistently. However, many uncertainties remain in how physical processes such as star formation and supernova feedback, which act below the scale of individual particles or cells, are implemented in these simulations. The computational cost of a single state-of-the-art hydrodynamical simulation is extremely high. This cost severely limits the number of simulations that can be performed, restricting the freedom to explore different parameter choices or alternative assumptions within a particular model. The computational demands of hydrodynamical simulations are compounded in the case of stellar halo models, in which the stars of interest constitute only $\\sim1\\%$ of the total stellar mass of a Milky Way-like galaxy. Even resolving the accreted dwarf satellites in which a significant proportion of these halo stars may originate is close to the limit of current simulations of disc galaxy formation. To date, few hydrodynamical simulations have focused explicitly on the accreted stellar halo (recent examples include Bekki \\& Chiba 2001; Brook {et~al.} 2004; Abadi {et~al.} 2006 and Zolotov {et~al.} 2009).  In the wider context of simulating the `universal' population of galaxies in representative ($\\gtrsim100\\,\\rm{Mpc^{3}}$) cosmological volumes, these practical limitations of hydrodynamical simulations have motivated the development of a powerful and highly successful alternative, which combines two distinct modelling techniques: well-understood high-resolution N-body simulations of large-scale structure evolved self-consistently from \\lcdm{} initial conditions and fast, adaptable semi-analytic models of galaxy formation with very low computational cost per run (Kauffmann, Nusser \\&  Steinmetz 1997; Kauffmann {et~al.} 1999; Springel {et~al.} 2001; Hatton {et~al.} 2003; Kang {et~al.} 2005; Springel {et~al.} 2005; Bower {et~al.} 2006; Croton {et~al.} 2006; De~Lucia {et~al.} 2006). In this paper we describe a technique motivated by this approach which exploits computationally expensive, ultra-high-resolution N-body simulations of \\textit{individual} dark matter haloes by combining them with semi-analytic models of galaxy formation. Since our aim is to study the spatial and kinematic properties of stellar haloes formed through the tidal disruption of satellite galaxies, our technique goes beyond standard semi-analytic treatments.  The key feature of the method presented here is the dynamical association of stellar populations (predicted by the semi-analytic component of the model) with sets of \\textit{individual particles} in the N-body component. We will refer to this technique as `particle tagging'. We show how it can be applied by combining the Aquarius suite of six \\newt{high resolution} isolated $\\sim10^{12}\\,\\rm{M_{\\sun}}$ dark matter haloes (Springel {et~al.} 2008a,b) with the \\galform{} semi-analytic model (Cole {et~al.} 1994, 2000; Bower {et~al.} 2006). These simulations \\newt{can resolve structures down to $\\sim10^{6}\\rm{M_{\\sun}}$, comparable to the least massive dark halo hosts inferred for Milky Way satellites (e.g. Strigari {et~al.} 2007; Walker {et~al.} 2009)}.  Previous implementations of \\newt{the particle-tagging approach} \\movt{(White \\& Springel 2000; Diemand, Madau \\& Moore 2005; Moore {et~al.} 2006; Bullock \\& Johnston 2005; De~Lucia \\& Helmi 2008)} have so far \\newt{relied on cosmological simulations} severely limited by resolution \\newt{(Diemand {et~al.} 2005; De~Lucia \\& Helmi 2008) or \\newt{else} simplified \\newt{higher resolution} N-body models (Bullock \\& Johnston 2005)}. \\movt{In the present paper, \\newt{we apply this technique} as a postprocessing operation to a `fully cosmological' simulation, in which structures have grown \\textit{ab initio}, interacting with one another self-consistently. \\newt{The resolution of our simulations is sufficient to resolve stellar halo substructure in considerable detail.}}  With the aim of presenting our modelling approach and exploring some of the principal features of our simulated stellar haloes, we proceed as follows. In \\mnsec{sec:method_overview} we review the Aquarius simulations and their post-processing with the \\galform{} model, and in \\mnsec{sec:buildinghaloes} we describe our method for recovering the spatial distribution of stellar populations in the halo by tagging particles. \\newt{We calibrate our model by comparing the statistical properties of the surviving satellite population to observations; the focus of this paper is on the stellar halo, rather on than the properties of these satellites.} In \\mnsec{sec:results} we describe our model stellar haloes and compare their structural properties to observations of the Milky Way and M31. We also examine the assembly history of the stellar haloes in detail (\\mnsec{sec:assembly}) and explore the relationship between the haloes and the surviving satellite population. Finally, we summarise our results in \\mnsec{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  We have presented a technique for extracting information on the spatial and kinematic properties of galactic stellar haloes that combines a very high resolution fully cosmological \\lcdm{} simulation with a semi-analytic model of galaxy formation. We have applied this technique to six simulations of isolated dark matter haloes similar to or slightly less massive than that of the Milky Way, adopting a fiducial set of paramter values in the semi-analytic model \\galform{}. The structural properties of the surviving satellites have been used as a constraint on the assignment of stellar populations to dark matter. We found that this technique results in satellite populations and stellar haloes in broad agreement with observations of the Milky Way and M31, if allowance is made for differences in dark halo mass.  Our method of assigning stellar populations to dark matter particles is, of course, a highly simplified approach to modelling star formation and stellar dynamics. The nature of star formation in dwarf galaxy haloes remains largely uncertain. In future, observations of satellites interpreted alongside high-resolution hydrodynamical simulations will test the validity of approaches such as ours. As a further simplification, our models do not account for a likely additional contribution to the halo from scattered {\\em in situ} (disc) stars, although we expect this contribution to be minimal far from the bulge and the disc plane. The results outlined here therefore address the history, structure and stellar populations of the accreted halo component in isolation.  Our results can be summarised as follows:  \\begin{itemize} \\item Our six stellar haloes are predominantly built by satellite accretion events occurring between $1<z<3$. They span a range of assembly histories, from `smooth' growth (with a number of roughly equally massive progenitors accreted steadily over a Hubble time) to growth in one or two discrete events. \\item Stellar haloes in our model are typically built from fewer than 5 significant contributors. These significant objects have stellar masses comparable to the brightest classical dwarf spheroidals of the Milky Way; by contrast, fewer faint satellites contribute to the halo than are present in the surviving population. \\item Typically, the most massive halo contributor is accreted at a lookback time of between 7 and 11 Gyr ($z\\sim1.5-3$) and deposits tidal debris over a wide radial range, dominating the contribution at large radii. Stars stripped from progenitors accreted at even earlier times usually dominate closer to the centre of the halo. \\item A significant fraction of the stellar halo consists of stars stripped from individual \\textit{surviving} galaxies, contrary to expectations from previous studies (e.g. Bullock \\& Johnston 2005). It is the most recent (and significant) contributors that are likely to be identifiable as surviving bound cores. Such objects have typically lost $\\sim90\\%$ of their original stellar mass. \\item We find approximately power-law density profiles for the stellar haloes in the range $10<r<100$ kpc. Those haloes formed by a superposition of several comparably massive progenitors have slopes similar to those suggested for the Milky Way and M31 haloes, while those dominated by a disproportionally massive progenitor have steeper slopes. \\item Our haloes have strongly prolate distributions of stellar mass in their inner regions ($c/a\\sim0.3$), with one exception, where an oblate, disc-like structure dominates the inner $10-20$ kpc. \\item Haloes with several significant progenitors show little or no radial variation in their mean metallicity ($Z/\\mathrm{Z_{\\sun}}$) up to 200~kpc. Those \\newt{in which} a small number of progenitors dominate show stronger metallicity gradients over their full extent or sharp transitions between regions of different metallicity. The centres of these haloes are typically more enriched than their outer regions. \\item The stellar populations of the halo are likely to be chemically enriched to a level comparable to that of the bright surviving satellites, but to be as old as the more metal-poor surviving `ultra faint' galaxies. The very metal-poor tail of the halo distribution is dominated by contributions from a plethora of faint galaxies that are insignificant contributors to the halo overall.   \\end{itemize}"
    },
    "0910/0910.4607_arXiv.txt": {
        "abstract": "{Detection of X-rays from classical novae, both in outburst and post-outburst, provides unique and crucial information about the explosion mechanism. Soft X-rays reveal the hot white dwarf photosphere, whenever hydrogen (H) nuclear burning is still on and expanding envelope is transparent enough, whereas harder X-rays give information about the ejecta and/or the accretion flow in the reborn cataclysmic variable. The duration of the supersoft X-ray emission phase is related to the turn-off of the classical nova, i.e., of the H-burning on top of the white dwarf core. A review of X-ray observations is presented, with a special emphasis on the implications for the duration of post-outburst steady H-burning and its theoretical explanation. The particular case of recurrent novae (both the \"standard\" objects and the recently discovered ones) will also be reviewed, in terms of theoretical feasibility of short recurrence periods, as well as regarding implications for scenarios of type Ia supernovae. } ",
        "introduction": "White dwarfs in cataclysmic variables explode as classical novae, when they accrete H-rich matter from their main sequence companion with some particular combinations of mass-accretion rate and initial white dwarf mass and luminosity. Matter accumulates on top of the white dwarf until hydrogen ignites, under degenerate or semi-degenerate conditions.This leads to a thermonuclear runaway, because self adjustment of the envelope once nuclear burning starts is not possible.  As a consequence of the explosion, a fraction of the accreted envelope (or even the whole envelope plus some material dredged-up from the white dwarf core) is ejected, with mass typically in the range $\\sim10^{-6}-10^{-4}$ M$_\\odot$, velocity of 100's or 1000's of km/s, and luminosity up to around $\\sim 10^{4-5}$ L$_\\odot$ (Prialnik \\& Kovetz 1995, Starrfield et al. 1998, Jos\\'e \\& Hernanz 1998). Steady nuclear burning is expected to take place in the remaining (if existing) H-rich envelope on top of the white dwarf. In fact, multiwavelength observations of classical novae have shown that a phase of constant bolometric luminosity (with values close to the Eddington limit) occurs after maximum, during the early optical decline; as the envelope mass is depleted, the photospheric radius decreases and thus the effective temperature increases, leading to a hardening of the spectrum, from optical to ultraviolet, extreme ultraviolet and soft X-rays, in agreement with theoretical predictions (Starrfield 1989, MacDonald 1996, Krautter 2002). Therefore, novae are expected to emit soft X-rays (E$\\le$1 keV) when steady H-bur\\-ning at the base of the white dwarf (WD) envelope heats the WD photosphere to effective temperatures from a few $10^5$ K up to $10^6$ K (Kato \\& Hachisu 1994, Sala \\& Hernanz 2005a, Kato these proceedings), provided that the expanding envelope is not opaque to X-rays; the nova then behaves as a supersoft source, SSS. However, if no H-rich envelope is left, or if the steady H-burning phase ends-up before the ejecta becomes optically thin to soft X-rays, the SSS phase will be absent. But such very short duration steady H-burning phases correspond to very tiny H-rich envelopes left after the explosion, much smaller than predicted by hydrodynamical models of nova explosions.  From the observation of the soft X-ray emission, the duration of the turn-off of the nova explosion can be measured; some other properties, such as the mass and metallicitof the remnant H-burning envelope and the mass of the underlying WD can be constrained, by comparing observations with current post-outburst nova models of this  phase (Sala \\& Hernanz 2005a). Additional information about the decline of classical novae comes from UV observations (Shore et al. 1996, G\\'onzalez-Riestra et al. 1998): turn-off times derived from UV data were in agreement with the soft X-ray ones, when both data were available.  There are other mechanisms of X-ray emission in classical novae, not related to residual H-nuclear burning on top of the white dwarf. Internal shocks in the expanding envelope, as well as shocks between the ejecta and circumstellar matter, lead to the heating of the plasma and the ensuing emission of X-rays, mainly by thermal bremsstrahlung; the corresponding energy of the emitted photons is larger than the energy of supersoft X-rays from the hot white dwarf photosphere. Such hard X-ray emission has been detected early after the explosion in some novae, e.g., V838 Her 1991 (5 days after optical outburst, which is the earliest detection of X-rays from a classical nova during outburst, Lloyd et al. 1992). In other novae it was detected later, e.g., 16 months after outburst in V 351 Pup 1991 (Orio et al. 1996), and both early and late in V382 Vel 1999 (15 days and 6 months after outburst, Orio et al. 2001). In the particular case of the recurrent nova RS Oph, hard X-rays up to $\\sim 50$ keV were detected very early after its 2006 outburst (Sokoloski et al. 2006, Bode et al. 2006), revealing the interaction between the ejecta and the wind of the red giant companion (RS Oph did not explode in a cataclysmic variable, as classical novae do; see next section).  It has also been suggested that the Comptonisation of the gamma-rays produced in $^{22}$Na nuclear decay could be responsible of prompt hard X-rays emitted by novae (Livio et al. 1992); but this has been ruled out after detailed modelling of the gamma-ray spectra of classical novae: Compton down-scattered photons are harder (energies larger than 20-30 keV, because of photoelectric absorption) and emitted earlier than the observed hard X-rays. In fact, emission of comptonised gamma-rays has very short duration (less than 2 days) and occurs before the nova is discovered optically (G\\'omez-Gomar et al 1998, Hernanz et al. 1999). The recent eruption of the recurrent nova RS Oph in 2006 was a good example of an early hard X-ray emission (detected with Swift/BAT and RXTE, Bode et al. 2006, Sokoloski et al. 2006, see above) non attributable to such Comptonisation of gamma-rays (Hernanz \\& Jos\\'e 2008a).  Finally, post-outburst novae also emit X-rays once accretion is reestablished, behaving then like cataclysmic variables (see for instance the case of V2487 Oph 1998, Hernanz \\& Sala 2002).  Hard X-ray emission itself is out of the scope of this review, but it should be mentioned that in some cases it is difficult to disentangle which fraction of the soft part of a whole X-ray spectrum corresponds to thermal plasma emission and which one to hot white dwarf photospheric emission. Only when ``canonical\" supersoft emission  is clearly seen, with large luminosities close to the Eddington limit and reasonable effective temperatures (up to about 100 eV), can it be ascribed to the whole hot white dwarf photosphere. ",
        "conclusions": "\\begin{itemize} \\item The number of novae detected as SSS has increased considerably in the last years, thanks mainly to Swift, XMM-Newton and Chandra. Some puzzling results concerning the temporal behaviour have been discovered, such as short period oscillations, sudden emergence and/or disappearance of soft X-ray emission. More observations are still needed to establish a classification and understand such diversity.  \\item Grating spectra are very rich, showing a wealth of lines; blueshifts of the absorption lines are specially relevant, since they reveal expansion during the super soft X-ray emitting phase (see Ness contribution to these proceedings). Detailed emission models (e.g., realistic WD atmospheres accounting for expansion effects) are in progress (see Van Rossum contribution to these proceedings).  \\item The WD mass and the envelope chemical composition - mainly its hydrogen content - determine the duration of the supersoft X-ray phase.  \\item Observed duration of the supersoft X-ray phase indicates, in general, M$_{\\rm env}$ $<$ M$_{\\rm accreted}$ - M$_{\\rm ejected}$ from hydrodynamical models. The mechanism leading to such extra mass loss is not completely clear. Wind mass loss clearly contributes (see Kato \\& Hachisu 1994 and subsequent papers). However, up to now the evolutionary hydrodynamical models that follow the whole nova outburst up to the mass-loss phase, do not obtain the tiny M$_{\\rm env}$ required, except for extreme combinations of accretion rate and initial white dwarf mass and luminosity (Yaron et al. 2005).  \\item Recurrent novae have an extremely short supersoft X-ray emitting phase, compatible with very small envelope mass left. These objects - like RS Oph - are very challenging for theory, since they are only possible for a narrow parameter range (extremely large WD mass and quite large accretion rate). They are good scenarios for type Ia supernovae, since WD mass is expected to increase, but a caveat is that massive WDs are made of ONe and not of CO, and ONe WDs do not explode but rather collapse once they reach the Chandrasekhar mass (Guti\\'errez et al. 1996).  \\end{itemize}  \\begin{figure} \\includegraphics[width=6cm,angle=-90]{mos12pnfit_onewd+mekal.ps} \\caption{Spectra of V5115 Sgr 2005a obtained with the EPIC (MOS1, MOS2 an pn) XMM-Newton cameras: normalised counts and residuals, for a ONe white dwarf atmosphere plus a thermal plasma best fit model. Model parameters: atmosphere has T$= 6\\times 10^5$K and L$_{\\rm WD}= 10^{36}$ erg/s, thermal plasma kT=2.5 keV and emission measure EM=$4 \\times 10^{55}$ cm$^{-3}$, photoelectric absorption N$_{\\rm H}=2 \\times 10^{21}$ cm$^{-2}$. The adopted distance is 10 kpc and the total unabsorbed luminosity is L$\\sim 10^{34}$ erg/s.} \\label{sgr05a} \\end{figure}  \\begin{table} \\centering% \\caption{Properties of some hydrodynamic models of nova explosions from Jos\\'e \\& Hernanz (1998) (see text for details).} \\label{tlab} \\begin{tabular}{lccc}\\hline Parameter                                   & ONe2 & ONe3 & ONe4\\\\ \\hline M$_{\\rm wd}(M_\\odot)$                       & 1.15   & 1.15    & 1.15\\\\ Mixing \\%                                   & 25     & 50      &  75\\\\ X (H mass fraction)                         & 0.53   & 0.35    & 0.18\\\\ Z (metals mass fraction)                    & 0.27   & 0.50    & 0.74\\\\ $\\Delta \\rm M_{\\rm env} (10^{-5}M_\\odot)$   & 3.2    & 3.2     & 3.5\\\\ $\\Delta \\rm M_{\\rm ejec} (10^{-5}M_\\odot)$  & 2.3    & 1.9     & 2.6\\\\ $\\Delta \\rm M_{\\rm rem} (10^{-5}M_\\odot)$   & 0.9    & 1.3     & 0.9\\\\ L$_{\\rm plateau} (10^4 L_\\odot)$            & 4.3    & 4.9     & 5.4\\\\ $\\tau_{\\rm nuc}$ (yr)                       & 8.9    & 7.4     & 2.4\\\\ \\hline \\end{tabular} \\end{table}"
    },
    "0910/0910.1214_arXiv.txt": {
        "abstract": "In addition to the maximum likelihood approach, there are two other methods which are commonly used to reconstruct the true redshift distribution from photometric redshift datasets: one uses a deconvolution method, and the other a convolution. We show how these two techniques are related, and how this relationship can be extended to include the study of galaxy scaling relations in photometric datasets. We then show what additional information photometric redshift algorithms must output so that they too can be used to study galaxy scaling relations, rather than just redshift distributions. We also argue that the convolution based approach may permit a more efficient selection of the objects for which calibration spectra are required. ",
        "introduction": "The next generation of sky surveys will provide reasonably accurate photometric redshift estimates, so there is considerable interest in the development of techniques which can use these noisy distance estimates to provide unbiased estimates of galaxy scaling relations. While there exist a number of methods for estimating photometric redshifts (Budavari 2009 and references therein), there are fewer for using these to estimate accurate redshift distributions (Padmanabhan et al. 2005; Sheth 2007; Lima et al. 2008; Cunha et al. 2009), the luminosity function (Sheth 2007), or the joint luminosity-size, color-magnitude, etc. relations (Rossi \\& Sheth 2008; Christlein et al. 2009; Rossi et al. 2010).  Ideally, the output from a photometric redshift estimator is a normalized likelihood function which gives the probability that the true redshift is $z$ given the observed colors (i.e. Bolzonella et al. 2000; Collister \\& Lahav 2004; Cunha et al. 2009). Let ${\\cal L}(z|{\\bm c})$ denote this quantity; it may be skewed, bimodal, or more generally it may assume any arbitrary shape.  Let $\\zeta$ denote the mean or the most probable value of this distribution (it does not matter which, although some of the logic which follows is more transparent if $\\zeta$ denotes the mean). Often, $\\zeta$ (sometimes with an estimate of the uncertainty on its value) is the only quantity which is available. Therefore, in Section~\\ref{dndz} we first consider how $\\zeta$ compares with the true redshift $z$, and contrast the convolution and deconvolution methods for estimating ${\\rm d}N/{\\rm d}z$ -- while in Section \\ref{cfc} we describe how to reconstruct the redshift distribution directly from colors. Section~\\ref{pdf} shows what this implies if one wishes to use the full distribution ${\\cal L}(z|{\\bm c})$. Section~\\ref{phil} shows how to extend the logic to the luminosity function, and Section~\\ref{phix} to scaling relations, again by contrasting the convolution and deconvolution methods, and showing what generalization of ${\\cal L}(z|{\\bm c})$ is required from the photometric redshift codes if one wishes to do this. A final section summarizes our results.  Where necessary, we write the Hubble constant as $H_0 = 100h~{\\rm km~s}^{-1}~{\\rm Mpc}^{-1}$, and we assume a spatially flat cosmological model with $(\\Omega_M,\\Omega_{\\Lambda}, h)=(0.3, 0.7, 0.7)$, where $\\Omega_M$ and $\\Omega_{\\Lambda}$ are the present-day densities of matter and cosmological constant scaled to the critical density. ",
        "conclusions": "We showed how previous work on deconvolution algorithms for making unbiased reconstructions of galaxy distributions and scaling relations (Sheth 2007; Rossi \\& Sheth 2008; Rossi et al. 2010) could be related to convolution-based methods. Whereas deconvolution based methods require accurate knowledge of $p(\\zeta|z)$, the distribution of the photometric redshift $\\zeta$ given the true redshift $z$, convolution based methods require accurate knowledge of $p(z|\\zeta)$.  Since $\\zeta$ is derived from photometry, this may more generally be written as $p(z|{\\bm c})$, where ${\\bm c}$ is the vector of observed photometric parameters which were used to estimate the redshift.  In both cases, $p(z|{\\bm c})$ and $p(\\zeta|z)$ are calibrated from a sample in which $z$ is known, and are then used in a larger sample where $z$ is not available.  If the smaller training set has the same selection limits as the larger dataset (e.g., both have the same magnitude limit) then both approaches are valid. \tWe illustrated our arguments with measurements in the SDSS (Figures~\\ref{pzzeta}--\\ref{NMconv}).  We also showed what additional information must be output from photometric redshift codes if their results are to be used in a convolution-like approach to provide unbiased estimates of galaxy scaling relations.  In particular, we argued that only if the redshift distribution output by a photo-$z$ algorithm, ${\\cal L}(z|{\\bm c})$, has the same shape as $p(z|{\\bm c})$, can the algorithm be said to be unbiased.  Only in this case its output (available for the full sample) can be used in place of $p(z|{\\bm c})$ (which is typically available for a small subset).  The safest way to accomplish this is for the training set to be a random subsample of the full dataset -- and to then tune the algorithm so that ${\\cal L}(z|{\\bm c}) = p(z|{\\bm c})$. If the training set is not representative, then care must be taken to ensure that ${\\cal L}(z|{\\bm c})$ does not yield biased results.  Obtaining spectra is expensive, so the question arises as to whether or not there is a more efficient alternative to the random sample approach.  For the convolution method, which requires $p(z|{\\bm c})$, the answer is clearly `yes'.  This is because some color combinations (e.g. the red sequence) might give rise to a narrow $p(z|{\\bm c})$ distribution, whereas others may result in broader distributions.  Since it will take fewer objects to accurately estimate the shape of a narrow $p(z|{\\bm c})$ distribution than a broad one, observational effort would be better placed in obtaining spectra for those objects which produce broad $p(z|{\\bm c})$ distributions. For the deconvolution approach, one would like to preferentially target those redshifts $z$ which produce broader $p(\\zeta|z)$ distributions -- for similar reasons.  But, since $z$ is not known until the spectra are taken, this cannot be done, so taking a random sample of the full dataset is the safest way to proceed.  Our methods permit accurate measurement of many scaling relations for which spectra were previously thought to be necessary (e.g. the color-magnitude relation, the size-surface brightness relation, the Photometric Fundamental Plane), so we hope that our work will permit photometric redshift surveys to provide more stringent constraints on galaxy formation models at a fraction of the cost of spectroscopic surveys."
    },
    "0910/0910.3950_arXiv.txt": {
        "abstract": "{The Constrained Minimal Supersymmetric Standard Model (CMSSM) is one of the simplest and most widely-studied supersymmetric extensions to the standard model of particle physics.  Nevertheless, current data do not sufficiently constrain the model parameters in a way completely independent of priors, statistical measures and scanning techniques.  We present a new technique for scanning supersymmetric parameter spaces, optimised for frequentist profile likelihood analyses and based on Genetic Algorithms.  We apply this technique to the CMSSM, taking into account existing collider and cosmological data in our global fit.  We compare our method to the \\MN~algorithm, an efficient Bayesian technique, paying particular attention to the best-fit points and implications for particle masses at the LHC and dark matter searches.  Our global best-fit point lies in the focus point region.  We find many high-likelihood points in both the stau co-annihilation and focus point regions, including a previously neglected section of the co-annihilation region at large $\\mzero$.  We show that there are many high-likelihood points in the CMSSM parameter space commonly missed by existing scanning techniques, especially at high masses.  This has a significant influence on the derived confidence regions for parameters and observables, and can dramatically change the entire statistical inference of such scans.} ",
        "introduction": "\\label{sec:intro}  New physics beyond the Standard Model (SM) is broadly conjectured to appear at TeV energy scales.  Particular attention has been paid to supersymmetric (SUSY) extensions of the SM, widely hoped to show up at the Large Hadron Collider (LHC).  One of the strongest motivations for physics at the new scale is the absence of any SM mechanism for protecting the Higgs mass against radiative corrections; this is known as the hierarchy or fine-tuning problem~\\cite{Martin:9709356}.  Softly-broken weak-scale supersymmetry (for an introduction, see Ref.~\\citealp{EWSUSY}) provides a natural solution to this problem via the cancellation of quadratic divergences in radiative corrections to the Higgs mass.  The natural connection between SUSY and grand unified theories (GUTs) also offers extensive scope for achieving gauge-coupling unification in this framework~\\cite{Ellis:1990}.  Supersymmetry has even turned out to be a natural component of many string theories, so it may be worth incorporating into extensions of the SM anyway (though in these models it is not at all necessary for SUSY to be detectable at low energies).  Another major theoretical motivation for supersymmetry is that most weak-scale versions contain a viable dark matter (DM) candidate~\\cite{DMreview}.  Its stability is typically achieved via a conserved discrete symmetry ($R$-parity) which arises naturally in some GUTs, and makes the lightest supersymmetric particle (LSP) stable.  Its `darkness' is achieved by having the LSP be a neutral particle, such as the lightest neutralino or sneutrino.  These are both weakly-interacting massive particles (WIMPs), making them prime dark matter candidates~\\cite{DMreview}.  The sneutrino is strongly constrained due to its large nuclear-scattering cross-section, but the neutralino remains arguably the leading candidate for DM.  Describing the low-energy behaviour of a supersymmetric model typically requires adding many new parameters to the SM.  This makes phenomenological analyses highly complicated.  Even upgrading the SM to its most minimal SUSY form, the Minimal Supersymmetric Standard Model (MSSM; for a recent review, see Ref.~\\citealp{Chung:0312378}), introduces more than a hundred free parameters.  All but one of these come from the soft terms in the SUSY-breaking sector.  Fortunately, extensive regions of the full MSSM parameter space are ruled out phenomenologically, as generic values of many of the new parameters allow flavour changing neutral currents (FCNCs) or CP violation at levels excluded by experiment.  One might be able to relate many of these seemingly-free parameters theoretically, dramatically reducing their number.  This requires specification of either the underlying SUSY-breaking mechanism itself, or a mediation mechanism by which SUSY-breaking would be conducted from some undiscovered particle sector to the known particle spectrum and its SUSY counterparts.  Several mediation mechanisms (for recent reviews, see e.g. Refs.~\\citealp{Chung:0312378} and~\\citealp{Luty:0509029}) have been proposed which relate the MSSM parameters in very different ways, but so far no clear preference has been established for one mechanism over another.  For a comparison of some mediation models using current data, see Ref.~\\citealp{AbdusSalam:09060957}.  Gravity-mediated SUSY breaking, based on supergravity unification, naturally leads to the suppression of many of the dangerous FCNC and CP-violating terms.  Its simplest version is known as minimal supergravity (mSUGRA)~\\cite{mSUGRA1,mSUGRA2}.  An alternative approach is to directly specify a phenomenological MSSM reduction at low energy.  Here one sets troublesome CP-violating and FCNC-generating terms to zero by hand, and further reduces the number of parameters by assuming high degrees of symmetry in e.g.~mass and mixing matrices.  A hybrid approach is to construct a phenomenological GUT-scale model, broadly motivated by the connection between SUSY and GUTs.  Here one imposes boundary conditions at the GUT scale ($\\sim$10$^{16}$\\,GeV) and then explores the low-energy phenomenology by means of the Renormalisation Group Equations (RGEs).  One of the most popular schemes is the Constrained MSSM (CMSSM)~\\cite{CMSSM}, which incorporates the phenomenologically-interesting parts of mSUGRA.  The CMSSM includes four continuous parameters: the ratio of the two Higgs vacuum expectation values ($\\tanb$), and the GUT-scale values of the SUSY-breaking scalar, gaugino and trilinear mass parameters ($\\mzero$, $\\mhalf$ and $\\azero$).  The sign of $\\mu$ (the MSSM Higgs/higgsino mass parameter) makes for one additional discrete parameter; its magnitude is determined by requiring that SUSY-breaking induces electroweak symmetry-breaking.  Despite greatly curbing the range of possible phenomenological consequences, the small number of parameters in the CMSSM has made it a tractable way to explore basic low-energy SUSY phenomenology.  Before drawing conclusions about model selection or parameter values from experimental data, one must choose a statistical framework to work in.  There are two very different fundamental interpretations of probability, resulting in two approaches to statistics (for a detailed discussion, see e.g. Ref.~\\citealp{Cowan:1998}).  Frequentism deals with relative frequencies, interpreting probability as the fraction of times an outcome would occur if a measurement were repeated an infinite number of times.  Bayesianism deals with subjective probabilities assigned to different hypotheses, interpreting probability as a measure of the degree of belief that a hypothesis is true.  The former is used for assigning statistical errors to measurements, whilst the latter can be used to quantify both statistical and systematic uncertainties.  In the Bayesian approach one is interested in the probability of a set of model parameters given some data, whereas in the frequentist approach the only quantity one can discuss is the probability of some dataset given a specific set of model parameters, i.e. a likelihood function.  In a frequentist framework, one simply maps a model's likelihood as a function of the model parameters.  The point with the highest likelihood is the best fit, and uncertainties upon the parameter values can be given by e.g. iso-likelihood contours in the model parameter space.  To obtain joint confidence intervals on a subset of parameters, the full likelihood is reduced to a lower-dimensional function by maximising it along the unwanted directions in the parameter space.  This is the profile likelihood procedure~\\cite[and references therein]{profilelike}.  In the Bayesian picture~\\cite{BayesRevs}, probabilities are directly assigned to different volumes in the parameter space.  One must therefore also consider the state of subjective knowledge about the relative probabilities of different parameter values, independent of the actual data; this is a prior.  In this case, the statistical measure is not the likelihood itself, but a prior-weighted likelihood known as the posterior probability density function (PDF).  Because this posterior is nothing but the joint PDF of all the parameters, constraints on a subset of model parameters are obtained by marginalising (i.e.~integrating) it over the unwanted parameters.  This marginalised posterior is then the Bayesian counterpart to the profile likelihood.  Bayesians report the posterior mean as the most-favoured point (given by the expectation values of the parameters according to the marginalised posterior), with uncertainties defined by surfaces containing set percentages of the total marginalised posterior, or `probability mass'.  One practically interesting consequence of including priors is that they are a powerful tool for estimating how robust a fit is.  If the posterior is strongly dependent on the prior, the data are not sufficient to constrain the model parameters.  It has been shown that the prior still plays a large role in Bayesian inferences in the CMSSM~\\cite{Trotta:08093792}.  If an actual detection occurs at the LHC, this dependency should disappear~\\cite{Roszkowski:09070594}.  Clearly the results of frequentist and Bayesian inferences will not coincide in general.  This is especially true if the model likelihood has a complex dependence on the parameters (i.e. not just a simple Gaussian form), and if insufficient experimental data is available.  Note that this is true even if the prior is taken to be flat; a flat prior in one parameter basis is certainly not flat in every such basis.  In the case of large sample limits both approaches give similar results, because the likelihood and posterior PDF both become almost Gaussian; this is why both are commonly used in scientific data analysis.  The first CMSSM parameter scans were performed on fixed grids in parameter space~\\cite{Drees:9207234,grid-cmssm}.  Predictions of e.g.~the relic density of the neutralino as a cold dark matter candidate or the Higgs/superpartner masses were computed for each point on the grid, and compared with experimental data.  In these earliest papers, points for which the predicted quantities were within an arbitrary confidence level (e.g.~1$\\sigma$, 2$\\sigma$) were deemed ``good''.  Because all accepted points are considered equivalently good, this method provides no way to determine points' relative goodnesses-of-fit, and precludes any deeper statistical interpretation of results.  The first attempts at statistical interpretation were to perform (frequentist) $\\chi^2$ analyses with grid scans~\\cite{chisquare-cmssm}.  Limited random scans were also done in some of these cases.  Despite some advantages of grid scans, their main drawback is that the number of points sampled in an $N$-dimensional space with $k$ points for each parameter grows as $k^{N}$, making the method extremely inefficient.  This is even true for spaces of moderate dimension like the CMSSM.  The lack of efficiency is mainly due to the complexity and highly non-linear nature of the mapping of the CMSSM parameters to physical observables; many important features of the parameter space can be missed by not using a high enough grid resolution.  Another class of techniques that has become popular in SUSY analyses is based on more sophisticated scanning algorithms.  These are techniques designed around the Bayesian requirement that a probability surface be mapped in such a way that the density of the resultant points is proportional to the actual probability.  However, the points they return can also be used in frequentist analyses.  Foremost amongst these techniques is the Markov Chain Monte Carlo (MCMC) method~\\cite{Baltz:0407039,Allanach:0507283,Allanach:0601089,Austri:0602028,Allanach:0609295,Roszkowski:0611173,Roszkowski:07052012,Allanach:07050487,Trotta:0609126,Roszkowski:07070622,Allanach:08061184,Allanach:08061923,Buchmueller:08084128,Roszkowski:09031279,Martinez:09024715,Balazs:09065012,Belanger:09065048,Buchmueller:09075568,Bechtle:09072589,Lafaye:07093985}, which has also been widely used in other branches of science, in particular cosmological data analysis~\\cite{MCMCCosmo}.  The MCMC method provides a greatly improved scanning efficiency in comparison to traditional grid searches, scaling as $kN$ instead of $k^{N}$ for an $N$-dimensional parameter space.  More recently, the framework of nested sampling~\\cite{SkillingNS} has come to prominence, particularly via the publicly-available implementation \\MN~\\cite{MultiNest}.  A handful of recent papers have explored the CMSSM parameter space or its observables using this technique~\\cite{Trotta:08093792,Martinez:09024715,Feroz:08074512,Feroz:09032487,Trotta:09060366,Roszkowski:09070594,Scott:FermiLAT,Scott:09084082}, as well as the higher-dimensional spaces of the Constrained Next-to-MSSM (CNMSSM)~\\cite{LopezFogliani:09064911} and phenomenological MSSM (pMSSM)~\\cite{pMSSM}. \\MN~was also the technique of choice in the supersymmetry-breaking study of Ref.~\\citealp{AbdusSalam:09060957}. A CMSSM scan with \\MN~takes roughly a factor of $\\sim$200 less computational effort than a full MCMC scan, whilst results obtained with both algorithms are identical (up to numerical noise)~\\cite{Trotta:08093792}.  Besides improved computational efficiency, MCMCs and nested sampling offer other convenient features for both frequentist and Bayesian analyses.  In a fully-defined statistical framework, all significant sources of uncertainty can be included, including theoretical uncertainties and our imperfect knowledge of the relevant SM parameters.  These can be introduced as additional `nuisance' parameters in the scans, and resultant profile likelihoods and posterior PDFs profiled/marginalised over them.  In a similar way, one can profile and marginalise over all parameters at once to make straightforward statistical inferences about any arbitrary function of the model parameters, like neutralino annihilation fluxes~\\cite{Roszkowski:07070622,Martinez:09024715,Scott:09084082} or cross-sections~\\cite{Scott:FermiLAT}.  The profiling/marginalisation takes almost no additional computational effort: profiling simply requires finding the sample with the highest likelihood in a list, whereas marginalisation, given the design of MCMCs and nested sampling, merely requires tallying the number of samples in the list.  Finally, it is straightforward to take into account all priors when using these techniques for Bayesian analyses.  Although the prior-dependence of Bayesian inference can be useful for determining the robustness of a fit, it may be considered undesirable when trying to draw concrete conclusions from the fitting procedure.  This is because the prior is a subjective quantity, and most researchers intuitively prefer their conclusions not to depend on subjective assessments.  In this case, the natural preference would be to rely on a profile likelihood analysis rather than one based on the posterior PDF.  The question then becomes: how does one effectively and efficiently explore a parameter space like the CMSSM, with its many finely-tuned regions, in the context of the profile likelihood?  As a first attempt to answer this question, the profile likelihood of the CMSSM was recently mapped with MCMCs~\\cite{Allanach:07050487} and \\MN~\\cite{Trotta:08093792}.  Despite the improved efficiency of these Bayesian methods with respect to grid searches by several orders of magnitude, they are not optimised to look for isolated points with large likelihoods.  They are thus very likely to entirely miss high-likelihood regions occupying very tiny volumes in the parameter space.  Such regions might have a strong impact on the final results of the profile likelihood scan\\footnote{Missing fine-tuned regions could even modify the posterior PDF if they are numerous or large enough.}, since the profile likelihood is normalised to the best-fit point and places all regions, no matter how small, on an equal footing.  It appears that in the case of the CMSSM there are many such fine-tuned regions.  This is seen in e.g.~CMSSM profile likelihood maps with different \\MN~scanning priors~\\cite{Trotta:08093792}.  Given that the profile likelihood is independent of the prior by definition, these results demonstrate that many high-likelihood regions are missed when using a scanning algorithm optimised for Bayesian statistics.  In order to make valid statistical inferences in the context of the profile likelihood, the first (and perhaps most crucial) step is to correctly locate the best-fit points.  Setting confidence limits and describing other statistical characteristics of the parameter space makes sense only if this first step is performed correctly.  If one wishes to work confidently in a frequentist framework, some alternative scanning method is clearly needed.  The method should be optimised for calculating the profile likelihood, rather than the Bayesian evidence or posterior.  Even if the results obtained with such a method turn out to be consistent with those of MCMCs and nested sampling, the exercise would greatly increase the utility and trustworthiness of those techniques.  In this paper, we employ a particular class of optimisation techniques known as Genetic Algorithms (GAs) to scan the CMSSM parameter space, performing a global frequentist fit to current collider and cosmological data.  GAs and other evolutionary algorithms have not yet been widely used in high energy physics or cosmology; to our knowledge, their only prior use in SUSY phenomenology~\\cite{Allanach:0406277} has been for purposes completely different to ours\\footnote{See e.g. Refs.~\\citealp{GenExpHepRef,GenPheHepRef,Teodorescu:08040369} for their use in high energy physics, Ref.~\\citealp{GenCosRef}~for uses in cosmology and general relativity, and Ref.~\\citealp{GenNucRef}~for applications to nuclear physics.}.  There are two main reasons GAs should perform well in profile likelihood scans.  Firstly, the sole design purpose of GAs is to maximise or minimise a function.  This is exactly what is required by the profile likelihood; in the absence of any need to marginalise (i.e.~integrate) over dimensions in the parameter space, it is more important that a search algorithm finds the best-fit peaks than accurately maps the likelihood surface at lower elevations.  Secondly, the ability of GAs to probe global extrema excels most clearly over other techniques when the parameter space is very large, complex or poorly understood; this is precisely the situation for SUSY models.  We focus exclusively on the CMSSM as our test-bed model, but the algorithms could easily be employed in higher-dimensional SUSY parameter spaces without considerable change in the scanning efficiency (as they scale as $kN$ for an $N$-dimensional parameter space).  We compare our profile likelihood results directly with those of the \\MN~scanning algorithm.  This means that we also compare indirectly with MCMC scans, because \\MN~and MCMCs give essentially identical results~\\cite{Trotta:08093792}. We find that GAs uncover many better-fit points than previous \\MN~scans.  This paper proceeds as follows: in~\\sec{sec:analysis} we briefly review the parameters of the CMSSM, the predicted observables, constraints and bounds on them from collider and cosmological observations.  We then introduce GAs as our specific scanning technique of choice.  We present and discuss results in~\\sec{sec:results}, comparing those obtained with GAs to those produced by \\MN.  We include the best-fit points and the highest-likelihood regions of the parameter space, as well as the implications for particle discovery at the LHC and in direct and indirect dark matter searches.  We draw conclusions and comment on future prospects in~\\sec{sec:concl}. ",
        "conclusions": "\\label{sec:concl}  Constraining the parameter space of the MSSM using existing data is under no circumstances an easy or straightforward task.  Even in the case of the CMSSM, a highly simplified and economical version of the model, the present data are not sufficient to constrain the parameters in a way completely independent of computational and statistical techniques.  There have been several efforts to study properties and predictions of different versions of the MSSM.  Many recent activities in this field have used scanning methods optimised for calculating the Bayesian evidence and posterior PDF.  Those analyses have been highly successful in revealing the complex structure of SUSY models, demonstrating that some patience will be required before we can place any strong constraints on their parameters.  The same Bayesian scanning methods have also been employed for frequentist analyses of the problem, particularly in the framework of the profile likelihood.  These methods are not optimised for such frequentist analyses, so care should be taken in applying them to such tasks.  We have employed a completely new scanning algorithm in this paper, based on Genetic Algorithms (GAs).  We have shown GAs to be a powerful tool for frequentist approaches to the problem of scanning the CMSSM parameter space.  We compared the outcomes of GA scans directly with those of the state-of-the-art Bayesian algorithm \\MN, in the framework of the CMSSM.  For this comparison, we mostly considered \\MN~scans with flat priors, but kept in mind that e.g. logarithmic priors give rise to higher-likelihood points at low masses in the CMSSM parameter space;  we justified this choice of priors.  Our results are very promising and quite surprising.  We found many new high-likelihood CMSSM points, which have a strong impact on the final statistical conclusions of the study.  These not only influence considerably the inferred high-likelihood regions and confidence levels on the parameter values, but also indicate that the applicability of the conventional Bayesian scanning techniques is highly questionable in a frequentist context.  Although our initial motivation in using GAs was to gain a correct estimate of the likelihood at the global best-fit point, which is crucial in a profile likelihood analysis, we also realised that they can find many new and interesting points in almost all the relevant regions of parameter space.  These points strongly affect the inferred confidence regions around the best-fit point.  Even though we cannot be confident of exactly how completely our algorithm is really mapping these high-likelihood regions, it has certainly covered large parts of them better than any previous algorithm.  We think that by improving the different ingredients of GAs, such as the crossover and mutation schemes, this ability might even be enhanced further.  We largely employed the standard, simplest versions of the genetic operators in our analysis, as well as very typical genetic parameters.  These turned out to work sufficiently well for our purposes.  Although we believe that tuning the algorithm might produce even more interesting results, it is good news that satisfactory results can be produced even with a very generic version.  This likely means that one can apply the method to more complicated SUSY models without extensive fine-tuning.  One interesting outcome of our scan is that the global best-fit point is found to be located in the focus point region, with a likelihood significantly larger than the best-fit point in the stau co-annihilation region (which in turn actually still has a higher likelihood than the global best-fit value obtained with \\MN, even with logarithmic priors).  The focus point region is favoured in our analysis over the co-annihilation region, in contrast to findings from some recent MCMC studies~\\cite{Buchmueller:08084128,Buchmueller:09075568}, where the opposite was strongly claimed.  We also found a rather large part of the stau co-annihilation region, consistent with all experimental data, located at high $\\mzero$.  This part of the co-annihilation region seems to have been missed in other recent scans.  All these results show that, at least in our particular setup, high masses, corresponding either to the FP or the COA regions, are by no means disfavoured by current data (except perhaps direct detection of dark matter).  The discrepancy between this finding and those of some other authors that the FP is disfavoured might originate in the different scanning algorithms employed, or in the different physics and likelihood calculations performed in each analysis.  We have however shown, by comparing our results with others produced using exactly the same setup except for the scanning algorithm, that one should not be at all confident that all the relevant points for a frequentist analysis can be found by scanning techniques optimised for Bayesian statistics, such as nested sampling and MCMCs.  We have also calculated some of the quantities most interesting in searches for SUSY at the LHC, and in direct and indirect searches for dark matter.  We showed that GAs found much better points compared to \\MN~almost everywhere in the interesting mass ranges of the lightest Higgs boson, gluino and neutralino.  We confirmed previous conclusions that the LHC is in principle able to investigate a large fraction of the high-likelihood points in the CMSSM parameter space if it explores sparticle masses up to around $3\\tev$.  As far as the Higgs mass is concerned, there are many points with rather low masses that, although sitting just below the low mass limit given by LEP, are globally very well fit to the experimental data.  In the context of dark matter searches, we noticed that the global best-fit point and much of the surrounding $1\\sigma$ confidence level region at high cross-sections are actually already mostly ruled out by direct detection limits, if one assumes the standard halo model to be accurate.  We also argued that some of these points may be tested by upcoming indirect detection experiments, in particular the \\emph{Fermi} gamma-ray space telescope.  Finally, we realised that the high-likelihood stau co-annihilation region at large $\\mzero$ introduces a new allowed region in the combination of the neutralino mass and self-annihilation cross-section, which (to our knowledge) has not been observed previously.  We also compared our algorithm with \\MN~in terms of speed and convergence, and argued that GAs are no worse than \\MN~in this respect.  GAs have a large potential for parallelisation, reducing considerably the time required for a typical run.  This property, as well as the fact that the computational effort scales linearly (i.e. as $kN$ for an $N$-dimensional parameter space), also makes GAs an excellent method for the frequentist exploration of higher-dimensional SUSY parameter spaces.  Finally, perhaps the bottom line of the present work is that we once again see that even the CMSSM, despite its simplicity, possesses a highly complex and poorly-understood structure, with many small, fine-tuned regions.  This makes investigation of the model parameter space very difficult and still very challenging for modern statistical scanning techniques.  Although the method proposed in this paper seems to outperform the usual Bayesian techniques in a frequentist analysis, it is important to remember that it may by no means be the final word in this direction.  Dependence of the results on the chosen statistical framework, measure and method calls for caution in drawing strong conclusions based on such scans.  The situation will of course improve significantly with additional constraints provided by forthcoming data."
    },
    "0910/0910.1487_arXiv.txt": {
        "abstract": "We report the results of our intensive intranight optical monitoring of 8 optically bright `radio-intermediate quasars' (RIQs) having  flat or inverted radio spectra. The monitoring was carried out in {\\it R-}band on 25 nights during 2005-09. On each night only one RIQ was monitored for a minimum duration of $\\sim$ 4 hours (the average being 5.2 hours per night). Using the CCD as an N-star photometer, an intranight optical variability (INOV) detection threshold of $\\sim$ 1--2\\% was achieved for the densely sampled differential light curves (DLCs) derived from our data. These observations amount to a large increase over those reported hitherto for this rare and sparsely studied class of quasars which can, however, play an important role in understanding the link between the dominant varieties of powerful AGN, namely the radio-quiet quasars (RQQs), radio-loud quasars (RLQs)  and blazars. Despite the probable presence of relativistically boosted nuclear jets, inferred from their flat/inverted radio spectra, clear evidence for INOV in our extensive observations was detected only on one night. Also, flux variation between two consecutive nights was clearly seen for one of the RIQs. These results demonstrate that as a class, RIQs are much less extreme in nuclear activity compared to blazars. The availability in the literature of INOV data for another 2 RIQs conforming to our selection criteria allowed us to enlarge the sample to 10 RIQs (monitored on a total of 42 nights for a minimum duration of $\\sim 4$ hours per night). The absence of large amplitude INOV $(\\psi \\geqslant 3\\%)$ persists in this enlarged sample. This extensive database has enabled us to arrive at the first estimate for the INOV Duty Cycle (DC) of RIQs. The DC is found to be small ($\\sim$ 9\\%), increasing to $\\sim$ 14\\% if the two cases of `probable' INOV are included. The corresponding value is known to be $\\sim 60\\%$ for BL Lacs and $\\approx 15\\%$ for both RLQs and RQQs, if they too are monitored for $\\ga 4 - 6$ hours in each session. Our observations also provide information about the long-term optical variability (LTOV) of RIQs, which is found to be fairly common and reaches typical amplitudes of $\\approx$ 0.1-mag. The light curves of these RIQs are briefly discussed in the context of a theoretical framework proposed earlier for linking this rare kind of quasars to the much better studied dominant classes of quasars. ",
        "introduction": "\\vskip0.5in  Although sparsely studied so far, radio-intermediate quasars (RIQs) can serve as an important tool for probing the relationship between the radio-loud quasars (RLQs), blazars and radio-quiet quasars (RQQs, also termed as radio weak quasars). Ever since the discovery of quasars, the dichotomy in their radio loudness has been debated as an outstanding issue, since about $\\sim 10\\%$ of optically selected quasars are found to be stronger emitters in the radio band compared to the remaining population (Sandage 1965; Strittmatter et al. 1980; Kellermann et al. 1989, 1994; Stocke et al. 1992).  Following Kellermann et al. (1989) who made VLA observations of optically selected Palomar-Green Bright Quasar Survey (BQS) objects, radio loudness is usually quantified in terms of a parameter `$R$', the ratio of continuum flux density at radio (5 GHz) and optical (4400 \\AA) wavelengths. The histogram of $R$ found in their study shows a large peak at $R \\la 0.1-3$ and a second, weaker peak at $R \\sim 100-1000$. Accordingly, quasars falling in the range $3 < R < 100$ can be termed as ``Radio Intermediate Quasars (RIQs)'' (e.g., Miller, Rawlings \\& Saunders 1993; see also, Diamond-Stanic et al. 2009). Although it has been argued that RLQs are merely the tail of a broad distribution of quasar radio strengths and not a separate population (e.g., Cirasuolo et al. 2003) careful analysis of the largest available samples strongly indicates that there is indeed a bimodal distribution in $R$ (Ivezi\\'c et al. 2002, 2004) and that the radio loud fraction increases with rising optical luminosity but decreases with increasing redshift (e.g., Jiang et al. 2007; Rafter, Crenshaw \\& Wiita 2009).  The observed radio luminosity of a RIQ is typically well above the value dividing the Fanaroff-Riley type I (FRI) and type II (FRII) radio galaxies [log$P_{20cm}$ (W/Hz) $\\simeq 25.0$] (Falcke, Gopal-Krishna \\& Biermann 1995b), but probably not so once the possibility of relativistic flux boosting is taken into account (e.g., Wang et al. 2006). It has been long been suggested that RIQs are Doppler boosted counterparts of RQQs, such that the relativistic jet is closely aligned to the line of sight. This was mainly inferred from the radio versus [O III] line intensity diagram for optically selected quasars and also from the high brightness temperature of the radio emission from RIQs, as well as their usually flat or inverted radio spectra and substantial radio flux variability (Miller, Rawlings \\& Saunders 1993; Falcke, Gopal-Krishna \\& Biermann 1995b; Falcke, Patnaik \\& Sherwood 1996a; Xu et al. 1999; cf. Barvainis et al. 2005; Falcke et al. 2001).  Implicit in this paradigm is the presence of relativistic jets in RQQs. This premise has received considerable support from deep radio imaging, which has revealed faint kpc-scale radio structures in several RQQs (e.g., Kellermann et al. 1994; Blundell \\& Rawlings 2001; Leipski et al. 2006). In addition, Blundell, Beasly \\& Bicknell (2003) have reported evidence for a highly relativistic parsec scale radio jet in the RQQ PG1407+263. Independent support for the possibility of RQQs possessing relativistic jets has come from the observations of intranight optical micro-variability in RQQs (e.g., Gopal-Krishna et al. 2000, 2003; Gupta and Joshi 2005; Czerny et al. 2008) and in their weaker cousins, Seyfert 1 galaxies (Jang \\& Miller 1995, 1997; Carini, Noble \\& Miller 2003). At least a modest amount of relativistic beaming in the jets of radio weak quasars has also been inferred from the detection of compact cores on VLBI scales and radio variability on month-like to year-like timescales, indicating minimum brightness temperatures in a broad range from $10^8$--$10^{11}$ K  (e.g., Ulvestad, Antonucci \\& Barvainis 2005; Barvainis et al. 2005; Wang et al. 2006). However, there is at present little clarity about the nature of jets in RQQs {\\it vis-\\'a-vis} those in RLQs and blazars whose jets are widely attributed to spinning super massive black holes (SMBHs) (Blandford 1990 and references therein; Wilson \\& Colbert 1995), or more massive SMBHs accreting at smaller fraction of the Eddington rate (e.g., Boroson 2002). A number of authors have drawn an analogy between the radio emission of quasars and X-ray binaries and have attributed the main difference between RLQs and RQQs to different accretion modes changing the nature of their jets (e.g., K\\\"{o}rding, Jester \\& Fender 2006). Stellar mass black holes are known to slide into a state where radio emission is strongly suppressed or quenched (e.g., Corbel et al. 2001; Maccarone, Gallo \\& Fender 2003). In this general picture RIQs could well be relativistically beamed counterparts of the RQQ jets. In alternative schemes, the radio weakness of RQQs is attributed to a rapid deceleration of the jet on subparsec scale, caused perhaps by: strong interaction with the torus material (e.g., Falcke, Malkan \\& Biermann 1995a; Falcke, Gopal-Krishna \\& Biermann 1995b); mass loading with the debris of stars tidally disrupted by the SMBH (Gopal-Krishna, Mangalam \\& Wiita 2008); ejection velocities less than the escape velocity (Ghisellini, Haardt \\& Matt 2004); or a high photon density environment produced by the quasar  (Barvainis et al. 2005).  If indeed, RIQs are relativistically boosted counterparts of RQQs, observing the former provides an opportunity to better detect RQQ jets and compare their properties with the better studied, more powerful radio jets of RLQs. For instance, it would be interesting to enquire if the postulated relativistic beaming of RQQ jets manifests itself as blazar-like behaviour in RIQs. Indeed, this seems to be the case for the most prominent example of the RIQ class, namely III Zw 2 (see below). However, essentially no such information is currently available for RIQs as a class even though they, despite being numerically rare, are a potentially very important link between the major AGN classes powered by nuclear activity, namely RQQs, RLQs and blazars. The present study is a first attempt to bridge this gap in a statistically significant manner.  In terms of radio properties, RIQs form a distinct subclass. Their dominant feature is a radio core (compact on arcsecond scale) of a flat, variable radio spectrum (Falcke, Patnaik \\& Sherwood 1996a; Kukula et al. 1998; also, Kellermann et al. 1994; Ulvestad et al. 2005). The aforementioned view that they are relativistically boosted RQQ jets has found strong support from the discoveries of long-term, large amplitude ({\\bf up to a factor of 20}) radio variability of the RIQ III Zw 2 (B0007$+$106; Ter\\\"asranta et al. 1998) and of superluminal motion in its radio core (Brunthaler et al.\\ 2000, 2005), both being canonical attributes of blazars. Unfortunately, equivalent observational data are severely lacking for other known RIQs. It may further be emphasized that III Zw 2 ($z = 0.089$) has in fact been classified as Seyfert 1 (Arp 1968; Osterbrock 1977) hosted by a spiral galaxy (Hutchings \\& Campbell 1983; Taylor et al. 1996). Thus, an interesting question is which, if any, blazar-like characteristics are associated with RIQs of high optical luminosity ($ M_{B} < -23.5$), which are bona-fide QSOs and hence more likely to be associated with massive elliptical galaxies. Promising indications come from the recent finding that for the radio spectra of the RIQs found in the SDSS are typically flat or inverted (consistent with the postulated relativistic beaming of their jets) and, moreover, radio flux variability (near 1 GHz) on year-like time scales is actually comparably common for RIQs and RLQs (Wang et al. 2006). For both these reasons, RIQs appear to be promising candidates for INOV, albeit it is unclear if the mechanism responsible for long-term variations is also responsible for INOV.  In recent years, INOV has been increasingly recognized as a signature of blazar-like jets.  Extensive observations have shown that INOV with large amplitude (with variations, $\\psi$, exceeding 3\\%) occurs almost exclusively in blazars and, furthermore, the duty cycle (DC) of INOV in blazars is very high ($> 50\\%$), provided the monitoring duration exceeds about 4 hours (e.g., Carini 1990; Gopal-Krishna et al. 2003; Stalin et al. 2004a, 2005). To exploit this clue we have carried out an INOV survey of a fairly large and representative sample of 8 RIQs, further augmented by the INOV data available in the literature for another 2 RIQs that meet the selection criteria adopted for our sample of 8 RIQs. Thus, the enlarged sample consists of 10 optically bright and intrinsically luminous, core-dominated RIQs. {\\bf Cosnsiderable evidence exisits for relativistic beaming in these RIQs. Falcke, Mathew \\& Biermann (1995a) find J1336+1725 to be radio variable, but not J1701+5149. For J1259+3423 multi-epoch flux measurements are not known to us, hence its variablity status is unknown at present. The remaining 7 RIQs in our sample are all found to be radio variable (Wang et al. 2006, Falcke, Sharwood \\& Patnaik 1996). Thus, at least 8 of the 10 RIQs are known to shown radio variability. Furthermore, spectral information is available in the literature for 8 out of the 10 RIQs and each case the radio spectrum is found to be either flat or inverted (see Table 1). All this suggests that the radio jets in essentially all there RIQs are at least modestly Doppler boosted. }. ",
        "conclusions": "The results reported here for a well defined, representative sample of {\\bf flat/inverted spectrum} RIQs provide a large increase over the existing information, both in terms of sample size and observing time. Thus, they permit the first good estimate of the INOV characteristics of RIQs and allow their comparison with those of the major AGN classes that are widely separated in the degree of radio loudness. The INOV mechanism for the major AGN classes continues to be debated (Section 1). For the (jet-dominated) blazars, INOV is believed to be associated with irregularities in the non-thermal Doppler boosted jet flow, impacted by shocks (e.g., Blandford \\& K{\\\"o}nigl 1979; Miller, Carini \\& Goodrich 1989; Marscher 1996). In contrast, the instabilities or perturbations within the accretion disc might contribute very significantly to the INOV in the case of RQQs (e.g., Wiita et al. 1991; Mangalam \\& Wiita 1993), particularly since any contribution from the jet must be weak.  To put the question of INOV of RIQs in perspective one may recall the known INOV characteristics of blazars (e.g., Heidt \\& Wagner 1996; Dai et al. 2001; Romero et al. 2002; Xie et al. 2002; Sagar et al. 2004; Stalin et al. 2005) and RQQs (e.g., Gopal-Krishna et al. 1993, 2003; Jang \\& Miller 1995, 1997;  Stalin et al. 2005; Carini et al. 2007). The major findings of these studies are that blazars tend to vary more frequently and more strongly ($\\psi > 3\\%$) on intranight timescales, in stark contrast to RQQs and even to non-blazar RLQs, which show only mild INOV (i.e., $\\psi < 3\\%$; Stalin et al. 2004b). Moreover, for monitoring durations in excess of $\\sim 4$ hours, the INOV duty cycle is $\\ga 60\\%$ for BL Lacs (Carini 1990; Miller \\& Noble 1996; Romero et al. 2002; Gopal-Krishna et al. 2003; Stalin et al. 2005), mostly with large INOV amplitudes ($\\psi > 3\\%$; Gopal-Krishna et al. 2003) but the duty cycle is only $\\la 15\\%$ for RQQs (Jang \\& Miller 1995, 1997; de Diego et al. 1998; Romero et al. 1999; Gopal-Krishna et al. 2003; Stalin et al. 2004b), or even core dominated RLQs (Stalin et al. 2004b; Ram{\\'i}rez et al. 2009).  Thus, the key result from the present study is that the INOV duty cycle for RIQs is small ($\\sim$ 10\\%), and clearly not greater than that known for RQQs and non-blazar type RLQs.  Although exact values of computed DCs will depend upon the sample and analysis technique, the key difference in INOV frequency is between blazars and  both radio loud and radio quiet non-blazar classes (Stalin et al. 2004b, 2005) as also confirmed recently by Ram{\\'i}rez et al. (2009). Moreover, on no occasion (among the 42 nights) do we have the evidence for an INOV amplitude exceeding $3\\%$ level, even though each of the 10 RIQs was monitored for at least $\\sim$ 4 hours in every session (Sect.\\ 1; Table 2). Specifically, the monitoring durations in our programme ranged between 3.9 hours to 7.6 hours, with an average of 5.2 hours. As emphasized by Carini (1990), the  duration of monitoring in a given session can play a substantial role in determining the INOV status of an AGN, such that the probability of confirmed INOV detection in BL Lacs is found to increase from $50\\%$ to $80\\%$ if the duration of monitoring is raised from $\\sim$ 3 hours to $\\sim$ 8 hours. A similar dependence on monitoring duration has been noticed for RQQs, where the chances of INOV detection increases to $\\sim 24\\%$ if the monitoring is done for $\\sim$6-7 hours (de Diego et al. 1998; Carini et al. 2007). It is conceivable that the INOV duty cycle of $\\sim 9 - 14\\%$ estimated here for RIQs may go up marginally when longer monitoring sessions become possible, but this is unlikely to alter our main conclusions.  The simplest explanation for the difference in variability is to assume that all quasars do possess nuclear radio jets, with the differences in both observed INOV DCs and amplitudes explained by different Doppler boosting arising from both different viewing angles and different shock velocities (e.g., Gopal-Krishna et al. 2003). Another possible explanation for the low INOV DC would involve dilution of jet's optical emission by the expected accretion disc emission; however, this is unlikely to be important here, since all the RIQs in our sample lie at small redshifts ($z < 1.5 $; Table 1) whereas the bulk of the continuum emission of the disk would be expected to arise in the far-UV and would only move in the R-band for $z > 2$ (e.g., Bachev et al. 2005; Carini et al. 2007). Furthermore, since all but one of these RIQs have $M_B < -24.5$, the chance of the host galaxies' stellar continua diluting the INOV when compared to the other classes of QSOs is also very small.  Were the INOV properties of AGNs to depend primarily on the beamed radio emission, one would expect stronger and more frequent INOV for our sample of RIQs, as compared to that found for RQQs. Such an expectation would be in tune with the widely held notion that RIQs are Doppler boosted counterparts of RQQs, as inferred from their high brightness temperature ($T_{B}$), prominent radio cores and low extended fluxes at radio wavelengths, measured using milli-arcsecond resolution radio images available for a few RIQs (Section 1). The INOV data reported here do not accord well these broad expectations, since we have found the INOV of RIQs to be as mild as that known for RQQs.  Thus, the RIQ PG0007+106 (III Zw 2) remains  the only well established case for which a blazar-like, large and rapid flux variability has been detected (both at radio and optical wavelengths, Sect.\\ 1). In particular, the source was found by Jang $\\&$ Miller (1997) to vary by up to 0.1 mag on each the two nights they monitored it for 4-hour duration. Interestingly, the observed high level of activity would not have been anticipated, given that this object was found to show a very low optical polarization: $P_{op} = 0.28 \\pm 0.19\\%$, (Berriman et al. 1990). The RIQ III Zw 2 underwent a slow fading by $\\sim$0.8 mag between 1978 to 1981, after which it displayed rapid flares of increasing amplitude, brightening by 0.92 magnitude in 27 days during November 1985 and then fading by 0.95 mag in the following week (Pica et al. 1988). Such large variability, along with the observed INOV (see above) is clearly characteristic of OVV type blazars. In the longer term, its optical light curve has shown a four-fold variation in 25 years (Salvi et al. 2002) and, likewise, a 20-fold brightening within 4 years has been recorded at radio wavelengths (Aller et al. 1985). All these variability properties firmly place this object in the blazar category (Kukula et al. 1998; Aller et al. 1999; Ter\\\"asranta et al. 2005). In contrast, none of the 10 RIQs in our representative sample qualify for the blazar classification, on the basis of the currently available data summarized above.  The present observations are also useful for measuring any long-term optical variability (LTOV) occurring on month-like or longer timescales. For the vast majority of RIQs, which are represented by the present sample and which are hosted by luminous (very likely elliptical) galaxies, we find the optical variability to be very common on month/year-like time scales, with typical amplitude approaching 0.1-mag level in R-band. This RIQ specific result is in accord with the findings reported by Webb \\& Malkan (2000) for more common AGN types. For roughly half the AGNs they found optical variability amplitudes of 0.1 -- 0.2 mag (rms) on month-like time scale. A similar pattern is reported by Barvainis et al. (2005), based on their 10-epoch VLA radio variability survey of a large sample consisting of RQQs, RIQs and RLQs. They found no statistical difference between the degrees of radio variability displayed by these three major subclasses of quasars. Thus, both intranight and long-term optical variability properties, together with the other afore-cited evidence from the literature, demonstrate that as a class RIQs display much lower activity levels than do blazars, which would be rather surprising if they are a strongly relativistically beamed subset of RQQs, a possibility widely discussed in the literature.   The possibility that these RIQ sources possess intrinsically moderate radio jets and are not strongly Doppler boosted (in either radio or optical band) therefore must be considered.  In that case either roughly comparable amounts of jet fluctuations, or even disc dominated variability, could explain the similar INOV and LTOV properties of RIQs and RQQs. The low optical polarizations seen for RIQs in our sample (Table 1) also seem consistent with a lack of strong relativistic beaming, though the case of III Zw 2 discussed above calls for caution in reaching this inference. A key observational difference between these RIQs and the RQQs is the predominance of flat or inverted radio spectra for the former group, which is generally not the case for RQQs (e.g., Kukula et al. 1998) and which was argued to imply (at least somewhat) boosted jets in the RIQs.  One conceivable explanation then for the relative lack of INOV in RIQs would be that, because of bending of their jets on the innermost scales, their optically emitting inner portions are misdirected and hence concealed from us (despite the beaming) and only the more extended radio emitting outer parts of the jets happen to be pointed towards us. Verifying this alternative would require sensitive VLBI observations. However, an implication of this scenario is that one would expect to find some cases of blazar-like optical variability in radio-quiet AGN whose radio emitting (but not optical emitting) sections of the jet are misaligned from us. The few searches made so far for radio-quiet BL Lacs have either been negative (e.g., Stocke et al. 1990; Londish et al. 2007) or have produced a only a small fraction of candidate BL Lacs that have low, but not extremely low, upper limits to their radio fluxes (Collinge et al. 2005). Finally, it may be noted that despite the vast increase in the intranight monitoring data, as reported here, the total number of nights used for their monitoring is still modest (42 nights) when compared with those devoted to blazars and even RQQs. Since a few convincing cases of RIQs showing modest intranight or internight variability have nonetheless been found, it would be worthwhile to continue such observations in the search for examples of blazar-like strong INOV activity among RIQs, similar to that detected so far only for the nearest known RIQ, III Zw 2."
    },
    "0910/0910.1678_arXiv.txt": {
        "abstract": "{}{A new method to constrain the cosmological equation of state is proposed by using combined samples of gamma-ray bursts (GRBs) and supernovae (SNeIa). }{The Chevallier-Polarski-Linder parameterization is adopted for the equation of state in order to find out  a realistic approach to achieve the deceleration/acceleration transition phase of dark energy models.}{We find that GRBs, calibrated by SNeIa, could be  good distance indicators capable of discriminating between cosmological models and $\\Lambda$CDM model at high redshift.}{} ",
        "introduction": "From an observational viewpoint, one of the fundamental goals of cosmology is  to measure cosmological distances and then to build up a suitable and reliable cosmic distance ladder. This issue has recently become even more important due to the evident degeneracy of several dark energy models with $\\Lambda$CDM, despite the advent of the so-called Precision cosmology, \\citep{Ellis}.  In the last two  decades, a class of accurate standard candles, the supernovae type Ia (SNeIa), has been highly studied and the results obtained from the use of these objects led to the surprising discovery of the  acceleration of the cosmic Hubble flow  (for a review see Kowalski et al 2008). However these objects are poorly detectable at redshifts higher than $\\sim 1.5$, so we need distance indicators at higher redshifts in order to remove the  degeneration of dark energy models affecting  current cosmological models ($\\Lambda$CDM is a good approximation of the observed Universe, even though there is still no theoretical basis about the nature of its components, but the issue of global evolution is far from being addressed; for a comprehensive review see Copeland et al 2006). A possible solution could be found by adopting gamma ray bursts (GRBs) as distance indicators.  GRBs are  the most powerful explosions in the Universe: the most likely scenarios for their generation are the formation of  massive black holes or the  coalescence of binary stellar systems. These events are observed at considerable distances, so there have been several efforts to frame them into the standard of the cosmological distance ladder. In the literature, there are several models that  account for  GRB formation \\citep{Meszaros}. All  these scenarios involve a similar shock phenomenon: a \"fireball\",  possibly be supported by a further jet emission. However none of these models is intrinsically capable of integrating all the observable quantities.  Despite  the poor knowledge of the GRB mechanism, it seems that GRBs could be used as reliable  distance indicators. There exist several observational correlations  among the photometric and spectral properties of GRBs to support this possibility \\citep{Leandros, Ghirl2}. Nevertheless the origin of these spectroscopic and photometrical correlations is not known very well and there are several efforts to interpret the behavior of GRB features in a coherent way, by  relatively simple scenarios  \\citep{MG, ghisellini}. Succeeding in explain the mechanism that generates  GRBs is one of the objectives of  modern astrophysics and to clarify these observed correlations in this context would make GRBs  reliable distance indicators. A complete review of the existing luminosity relations for GRBs can be found in Schaefer (2007).  In this paper, we consider two relations, the one by Liang-Zhang (LZ) \\citep{LZ}, and the one by Ghirlanda (GGL) \\citep{GGL}. They are the only 3-parameter relations  known and  have less scatter with respect to the theoretical best fit than the other 2-parameter ones. However calibration of the relations used has been necessary in order to avoid the circularity problem. This means that all the relations need to be calibrated for each set of cosmological parameters. Indeed, all GRB distances, obtained in a photometric way, are strictly dependent on the cosmological parameters since, currently, there is no low-redshift ($z$ up to 0.2-0.3) set of GRBs available to achieve a cosmology-independent calibration. In order to overcome this difficulty, Liang et al. (2008), proposed a method in which several GRB relations have been calibrated by SNeIa. Supposing that these relations work at all redshifts and that, at the same redshift, GRBs and SNeIa have the same luminosity distance, it becomes possible, in principle, to calibrate the GRB relations using an interpolation algorithm. In this way,  it becomes possible to  build a GRB-Hubble diagram by calculating the luminosity distance for each GRB with the well-known relation between the luminosity distance $d_l$ and the energy-flux ratio of the distance indicators.  In the literature there are several paper that use similar methods to constraint the cosmological parameters of the Concordance Model using GRBs as extension of the SNeIa Hubble Diagram, \\citep{Firmani1,Firmani2,Wang1,Wang2,Li1,Li}.  Here we take into account a cosmological EoS working at any redshift, using GRBs as tracers and adopting again the Chevallier-Polarski-Linder  (CPL) parameterization. In particular we discuss a method which should allow us to obtain an analytic cosmology-independent formulation of the luminosity distance and then of the distance modulus. After a brief introduction to the GRB luminosity relations, we show fits of the data obtained by these relations and the results and perspectives of the approach are discussed in the last Section. ",
        "conclusions": "Starting from the Friedmann equation, we have investigated a new method to constrain the cosmological equation of state at high redshifts.  In particular we obtain an analytical formula for the distance modulus so we could directly estimate  the parameters of the cosmological model considered. The working hypothesis involves the use of GRBs as distance indicators  at  high redshift, well beyond the distance where SNeIa have been detected  to date.  The CPL parameterization for the EoS has been explicitly used for the whole matter-energy content of the Universe as a suitable approach to investigate the parameter $w=w(z)$ and discriminate values with respect to the $\\Lambda$CDM model. In particular, regarding the Friedmann equations, we have obtained, in the case of the LZ relation, the epoch for the transition between the deceleration-acceleration phases at a redshift value of $z \\approx 5 $,  with a reliable confidence level. This is a value that,  if higher than the redshift of the farthest GRB used, could be in agreement with current quasar formation scenarios. Also, we are in good agreement with the observed phantom/quintessence regime at the present epoch, that is for $z \\rightarrow 0$, we obtain $w \\leq -1$.  So we reject the current phantom regime by this analysis, obtaining for $w_0$ a value  in agreement with the $\\Lambda$CDM model at the present epoch. The method, while preliminary, seems to indicate that GRBs could be used as standard candles once a reliable unified model of their  photometric and spectroscopic quantities is achieved (some relevant results  are presented in Ghisellini et al 2008). However, more robust samples of data are needed and a more realistic EoS  (with respect to the simple perfect fluid models) should be taken into account in order to suitably track  redshift at any epoch (see for example Capozziello et al 2006).  With  improving observations, in particular with the launch of new satellites devoted to  GRB surveys, such as Fermi-GLAST\\footnote[3]{http://fermi.gsfc.nasa.gov} and AGILE\\footnote[4]{http://agile.rm.iasf.cnr.it}, one should be able to expand the samples of GRBs, possibly with data coming from objects at higher redshift.  Considering  these preliminary results, it seems that GRBs could be considered as a useful tool to remove degeneracy and constrain self-consistent cosmological models. The matching with other distance indicators  would improve the consistency of the Hubble distance-redshift  diagram by extending it up to redshift $6\\, -\\, 7$ and higher.           \\begin{table} \\caption{Results of the fits corrected for the 3 ``outying'' GRBs. SNeIa is  for the supernova Ia data, LZ is for the GRBs data obtained from the Liang-Zhang relation, GGL for the Ghirlanda et al. one.} % \\label{table:no3} % \\centering % \\begin{tabular}{l c c c} % \\hline\\hline % Relation & $w_0$ & $w_a$ & $R^2$ \\\\ % \\hline % LZ + SNeIa & $-0.95 \\pm 0.01$ & $0.74 \\pm 0.01$ & $0.999$ \\\\ GGL + SNeIa & $-0.865 \\pm 0.005$ & $0.66 \\pm 0.005$ & $0.999$ \\\\ \\hline % \\end{tabular} \\end{table}"
    },
    "0910/0910.1352_arXiv.txt": {
        "abstract": "We report on the results of HST optical and UV imaging, Spitzer mid-IR photometry and optical spectroscopy of a sample of 30 low-redshift ($z$$\\sim$0.1 to 0.3) galaxies chosen from the SDSS and GALEX surveys to be accurate local analogs of the high-redshift Lyman Break Galaxies (LBGs). The Lyman Break Analogs (LBAs) are similar in stellar mass, metallicity, dust extinction, SFR, physical size and gas velocity dispersion, thus enabling a detailed investigation of many processes that are important in star forming galaxies at high redshift. The main optical emission line properties of LBAs, including evidence for outflows, are also similar to those typically found at high redshift. This indicates that the conditions in their interstellar medium are comparable. In the UV, LBAs are characterized by complexes of massive clumps of star-formation, while in the optical they most often show evidence for (post-)mergers and interactions. In six cases, we find a single extremely massive (up to several $\\times 10^9 M_{\\odot}$) compact (radius $\\sim 10^2$ pc) dominant central object (DCO). The DCOs are preferentially found in LBAs with the highest mid-IR luminosities ($L_{24\\mu m} = 10^{10.3}$ to $10^{11.2}$ L$_{\\odot}$) and correspondingly high star formation rates (15 to 100 M$_{\\odot}$ per year). We show that the massive star-forming clumps (including the DCOs) have masses much larger than the nuclear super star clusters seen in normal late type galaxies. However, the DCOs do have masses, sizes, and densities similar to the excess-light/central-cusps seen in typical elliptical galaxies with masses similar to the LBA galaxies. We suggest that the DCOs form in the present-day examples of the dissipative mergers at high-redshift that are believed to have produced the central-cusps in local ellipticals (consistent with the disturbed optical morphologies of the LBAs). More generally, the properties of the LBAs are consistent with the idea that instabilities in a gas-rich disk lead to very massive star-forming clumps that eventually coalesce to form a spheroid.  Finally we comment on the apparent lack of energetically significant AGN in the DCOs. We speculate that the DCOs are too young at present to be growing a supermassive black hole because they are still in a supernova-dominated outflow phase (age less than 50 Myr). ",
        "introduction": "One of the most remarkable features of high redshift ($z\\sim2-6$) galaxies is their great potential for vigorous star formation. Moreover, observations with the {\\it Hubble Space Telescope} (HST) indicate that that star formation largely occurs in an extremely ``clumpy'' and compact mode \\citep[e.g.][]{cowie95,franx97,elmegreen05,stark08}. These clumps are reminiscent of star-forming complexes in local \\hii\\ regions, but their total masses and scale sizes are several orders of magnitudes larger. As shown by \\citet{elmegreen09}, clumps with masses of $10^{7-9}$ $M_\\odot$ and sizes of $\\gtrsim$1 kpc are a very typical constituent of extended, irregular galaxies at $z=1-4$ in the Ultra Deep Field \\citep[UDF,][]{beckwith06}. The relatively high gas velocity dispersions observed in some high redshift galaxies \\citep[e.g.][]{law07,law09,forster09} coupled with their large inferred gas column densities and gas fractions can explain the formation of the clumps and their bulk properties from gravitational instabilities in gas-rich, Jeans-unstable disks. This process may be instigated, for example, by the continuous accretion of gas or by the delivery of gas-rich material in small merger events \\citep{bournaud09,elmegreen09}. In any case, the compactness of the stellar clumps furthermore implies that they can form relatively unhampered by velocity shear or supernova feedback. Simulations show that for some galaxies the coalescence of clump systems may lie at the root of the formation of galaxy bulges, and possibly even their supermassive nuclear black holes if it is assumed that each clump gave rise to its own intermediate mass black hole prior to their migration to the center \\citep{noguchi99,immeli04,elmegreen08a,elmegreen08b}.  It would be extremely valuable to study similar objects in nearby galaxies where multi-waveband data with relatively high spatial resolution and sensitivity would allow the relevant physical processes to be probed in much greater detail. While local starburst galaxies used in previous studies have some similarities to star-forming galaxies at high redshift, there are some important differences: local starbursts having SFRs of $\\gtrsim$100 $M_\\odot$ yr$^{-1}$ are very dusty systems (the so-called Ultra-Luminous Infrared Galaxies, ULIRGs). The intense star formation occurs in compact, sub-kpc scale regions in the centers of interacting galaxies, and are usually inconspicuous in the far-UV \\citep{heckman98,goldader02}. Although very dusty starbursts clearly make up a non-negligible fraction of high redshift galaxies \\citep[e.g.][]{hughes98,huang05,papovich06,caputi07,burgarella09,daddi09}, most of the systems contributing to the cosmic SFR at high redshift are accounted for in deep UV-selected samples of Lyman Break Galaxies (LBGs) \\citep[e.g.][]{reddy06,reddy09,bouwens09}, indicating that the amount of dust extinction in LBGs is typically much lower: the average ratio of their bolometric IR-to-UV luminosities, $L_{IR}/L_{UV}$, is much lower ($\\sim10\\times$) compared to local galaxies of the same $L_{bol}$ \\citep[e.g.][]{buat07,reddy06,burgarella07}.  Locally calibrated dust corrections that are based on the reddening of the UV continuum slope as derived by \\citet{meurer99} seem to apply to typical LBGs \\citep{reddy06} with some exceptions \\citep[e.g.][]{baker01,reddy06,siana08}, and can be used to show that, despite their UV selection, LBGs at $z\\simeq2-4$ still put out most of their energy in the mid- and far-IR \\citep{bouwens09}.  The star formation in LBGs typically occurs over the entire extent of the galaxy ranging from a few hundred pc to several kpc \\citep[e.g.][]{franx97,giavalisco02,bouwens04,ferguson04,stark08}. The relatively low extinctions and high UV surface brightnesses seen in LBGs are characteristic of local blue compact dwarf galaxies (BCDs). However, typical BCDs have star formation rates up to two orders of magnitude smaller than LBGs and occur in galaxies with much lower mass \\citep[e.g.][]{telles97,hopkins02}, although some may resemble fainter LBGs at high redshift in their basic properties \\citep[see, e.g.,][]{ostlin09}. The main differences between LBGs and typical local starbursts studied previously may be due to a systematically lower metallicity and higher gas-mass fraction in the LBGs and/or due to a different triggering mechanism \\citep[e.g.][]{erb06a,erb06b,law09}.  The ``Lyman break analogs'' (LBA) project was designed in order to search for local starburst galaxies that share typical characteristics of high redshift LBGs \\citep{heckman05}. The LBA sample and main properties are given in \\citet{heckman05}, \\citet{hoopes07}, \\citet{basu-zych07,basu-zych09} and \\citet[][Paper I]{overzier08}, and we note that the LBAs are identical to the sample of ``supercompact UV-luminous galaxies'' (ScUVLGs) referred to in the papers listed above. In brief, the UV imaging survey performed by the Galaxy Evolution Explorer (GALEX) was used in order to select the most luminous ($L_{FUV}>10^{10.3}$ $L_\\odot$) and most compact ($I_{FUV}>10^{9}$ $L_\\odot$ kpc$^{-2}$) star-forming galaxies at $z<0.3$. Galaxies having these high $L_{FUV}$ are already extremely rare (local space density of $\\sim10^{-5}$ Mpc$^{-3}$), and the surface brightness requirement further reduces their absolute number density by a factor of $\\sim6$. These objects tend to be much more luminous in the UV than typical local starburst galaxies studied previously \\citep{heckman98,meurer97,meurer99} consistent with high SFRs and relatively little dust extinction. The median absolute UV magnitude of the sample is --20.3, corresponding to $\\simeq0.5L_{z=3}^*$, where $L_{z=3}^*$ is the characteristic luminosity of LBGs at $z\\sim3$ \\citep[][i.e., $M_{1700,AB}=-21.07$]{steidel99}. Analysis of their spectra from the Sloan Digital Sky Survey (SDSS) and their spectral energy distributions from SDSS and GALEX subsequently showed that the LBAs are similar to LBGs in their basic global properties, including: stellar mass, metallicity, dust extinction, SFR, physical size and gas velocity dispersion. In \\citet{overzier08} we showed that most of the UV emission in a preliminary sample of 8 LBAs originates in highly compact burst regions in small, clumpy galaxies that are also morphologically similar to LBGs. We demonstrated that {\\it if} LBGs at high redshift are also small merging galaxies similar to the LBAs, unfortunately this would be very hard to detect given the much poorer physical resolution and sensitivity. Furthermore, for the present paper it is important to point out that the LBAs also occupy an ``offset'' region in the main optical emission line diagnostics diagram (i.e. log(\\nii/\\ha) versus log(\\oiii/\\hb), or the ``BPT'' \\citep{baldwin81} diagram) analogous to galaxies at high redshift. This suggests that similar physical conditions may apply to their interstellar medium (ISM) as well. In summary, the LBAs appear indeed good local analogs to the LBGs.  In this paper we will present new results from our HST and Spitzer imaging campaigns and spectroscopic follow-up in order to highlight some of the remarkable properties of the LBA sample: we will compare the UV--optical sizes and morphologies, showing evidence for massive star-forming ``clumps'' that dominate the UV (\\S3). We explore possible connections between the starburst regions and the general properties of the optical emission line gas (\\S3.5 and \\S4), as well the relationship between the formation and evolution of massive stellar clumps in the context of the main structural components of young galaxies (\\S5). We conclude the paper with a summary and final remarks (\\S6). Forthcoming work on the LBA sample currently in preparation includes a detailed morphological analysis, multi-waveband modeling, mid-IR spectroscopy, VLBI radio imaging, X-ray spectroscopy, and integral field and long-slit optical and near-IR spectroscopy. ",
        "conclusions": "We have presented the results of the analysis of a sample of 30 low-redshift ($z <0.3$) galaxies selected from the union of the GALEX and SDSS surveys to have the high far-UV luminosities and small sizes characteristic of the high-redshift Lyman Break Galaxies \\citep{heckman05,hoopes07,basu-zych07,overzier08}. These ``Lyman Break Analogs'' (LBAs) offer the opportunity to investigate the physical processes occurring in Lyman Break Galaxies (LBGs) in considerably more detail than is possible at high-redshift\\footnote{\\citet{heckman05} coined the term ``living fossils'' in relation to the Lyman break analogs sample then discovered. Indeed, an apt analogy in the field of evolutionary biology is presented by the Coelacanth. A modern-day analog of this species of fish, known from the fossil record and believed to be extinct since the end of the Cretaceous period, was discovered in 1938. The find illustrates the use of analogs in the study of cosmology. Even though the local analogs will differ from their distant look-a-likes to some level of detail, and even though their local environment has probably changed drastically, they offer an entirely new way of investigating the fossil record alongside with other, more traditional lines of research.}. In this paper we reported on the internal structures of the LBA's using HST optical and UV images. We complemented these data with mid/far-IR photometry and optical spectroscopy from the SDSS and the VLT.  We have shown that the morphologies of the LBAs are dominated by massive compact clumps of young stars. {\\it Our most unexpected and potentially important result is that we have found six galaxies in which the galaxy UV morphology is dominated by a single very compact object which is unresolved or only marginally resolved in the HST images.} These dominant compact objects (``DCOs'') have estimated masses of few $\\times 10^8$ to few $\\times 10^9$ M$_{\\odot}$ and radii of-order 100 pc or less. In the other LBAs we instead find multiple clumps per galaxy, the brightest of which have estimated masses of typically a few $\\times 10^8$ $M_{\\odot}$ and radii of several hundred pc.  We show that the galaxies containing the DCOs differ from the other LBAs in many respects:\\\\  1. Their host galaxies tend to be more massive than the other LBAs (median masses of $\\sim 3 \\times 10^{10}$ vs. $\\sim 6 \\times 10^9$ $M_{\\odot}$ respectively).\\\\  2. They have systematically higher central concentrations of light than the other LBAs (the DCO is typically located near the center of a larger disk or envelope).\\\\  3. The DCOs are found in galaxies with high mid-infrared luminosities: L$_{24\\mu m} \\sim 10^{10.3}$ to $\\sim 10^{11.2}$ $L_{\\odot}$, roughly an order-of-magnitude more luminous than the other LBAs and comparable to luminous infrared galaxies.\\\\  4. The galaxies with a DCO are more strongly displaced from the locus of normal star forming galaxies in the standard BPT diagnostic emission-line ratio diagram. This displacement is in the sense of having larger \\oiii/H$\\beta$ and/or larger \\nii/H$\\alpha$ ratios.\\\\  5. The DCOs tend to occur in galaxies with higher gas densities and pressures as traced by the ionized gas ($P/k \\sim 10^7$ cm$^{-3}$ K).\\\\  6. The DCOs are found in galaxies in which the star-formation rate inferred from the dust-corrected H$\\alpha$ luminosity tends to be smaller than that derived from the mid-IR luminosity and/or the extinction-corrected far-UV luminosity. In the most extreme cases, this discrepancy reaches roughly an order-of-magnitude.\\\\  We have shown that the DCOs are not Type 1 (unobscured) AGN. We find no evidence for the presence of a very broad (several thousand km/s) H$\\alpha$ emission-line, with the most stringent constraints provided by our VLT GIRAFFE/ARGUS IFU spectra. Instead, these data show blue-asymmetric wings on the H$\\alpha$ and \\nii\\ line profiles indicative of a dusty outflow of ionized gas. While we can not entirely exclude the possibility that an energetically significant Type 2 (obscured) AGN is present in some or all of the DCOs, we regard this possibility as unlikely. In particular, the DCOs are associated with galaxies having unusually weak \\sii\\ and \\oi\\ emission-lines compared to normal star forming galaxies. This is the opposite of what is seen in starburst-AGN composite systems \\citep[e.g.][]{kewley06}, but similar to what is seen in low redshift IR-warm starbursts \\citep{kewley01}.  The continuum of the DCOs is dominated by the light of young stars: these are clearly strong highly compact starbursts. This is consistent with both the high gas pressures and emission-line profiles: both are commonly observed in starbursts and are the signposts of supernova-driven galactic winds \\citep[e.g.][]{lehnert96}. We have considered two specific scenarios under which the starbursts associated with the DCOs might be unusual. They may be evolved short-duration starbursts in which the rate of photoionization is relatively low because few O stars are present. They may also be ``leaky'' starbursts in which a significant fraction of the ionizing radiation escapes (matter bounded conditions). The latter scenario seems less likely, since it is difficult to envision circumstances in which most of the ionizing radiation escapes the galaxy, but most of the non-ionizing UV radiation is absorbed by dust and reprocessed in the mid/far IR.  We have compared the properties of the DCOs and other massive clumps to those of the nuclear star clusters typically seen in different kinds of nearby galaxies. The DCOs are highly massive and dense objects that do not appear to be related to single, young star clusters which are significantly less massive ($\\leq 10^7$ M$_{\\odot}$) and smaller ($\\sim$10 pc) than the structures we see in the LBAs. However, the DCO structures are quite similar to predictions for the merger-driven starbursts that are believed to be the progenitors of the extra light or central cusps seen generically in lower-mass ($\\leq M_*$) elliptical galaxies \\citep{kormendy08,hopkins09}. This would be consistent with our optical images of the LBAs, most of which strongly suggest on-going mergers. However, at present it is unclear whether the DCO hosts could passively evolve to form a full-grown elliptical without additional merging.  More generally, the properties of the LBAs are quite consistent with the idea that large Jeans instabilities in a gas-rich disk will lead to the formation of highly massive star forming clumps. At high redshift, it is not yet clear what the main source of this gas is: it could be replenished continuously, or come in the form of discrete accretion events of small galaxies or clouds (or, most likely, a combination thereof). These processes have been proposed to explain the properties of high-$z$ star-forming galaxies \\citep[e.g., see][and references therein]{overzier08,basu-zych09,elmegreen09,forster09,law07,law09}. In any event, it is expected that the massive stellar clumps will aggregate at the center of a galaxy via dynamical friction to ultimately form a bulge and supermassive black hole \\citep[][and references therein]{elmegreen09}.  The LBAs offer us the opportunity to study such processes in more detail, and compare samples across different redshifts.  We speculated about why we do not see clear evidence for the formation or rapid growth of a supermassive black hole in the DCOs (which appear to be the ideal black hole nurseries). One intriguing possibility is that these objects are simply too young for a black hole to be forming yet. Following \\citet{norman88} and \\citet{davies07}, the growth of a supermassive black hole inside a very massive compact star cluster may not commence until the cluster is old enough for the core-collapse supernovae phase to end (cluster age $>$ 50 Myr). Subsequent to this, gas shed by evolved stars will be retained in the cluster's potential well rather than being blown out. This gas can then cool, flow inward, and be used to grow the black hole. If correct, we should be able to identify a population of the descendants of the young DCOs in a phase as powerful AGN. This scenario might also account for the rarity of strong AGN in the LBG population \\citep[e.g.][]{steidel02,ouchixx}.  Follow-up UV spectroscopy using the HST Cosmic Origins Spectrograph has been allocated for selected targets in Cycle 17. The UV spectra will allow us to study many of the features discussed in this paper that are relevant to LBAs and to starbursts in general. These include the role of galactic outflows and low-luminosity AGN, the escape of ionizing radiation, the strength of Ly$\\alpha$\\ emission, stellar abundances and the initial mass function."
    },
    "0910/0910.0742_arXiv.txt": {
        "abstract": "We investigate constraints on the time variation of the fine structure constant between the recombination epoch and the present epoch, $\\Delta\\alpha/\\alpha \\equiv (\\alpha_{\\mathrm{rec}} - \\alpha_{\\mathrm{now}})/\\alpha_{\\mathrm{now}}$, from cosmic microwave background (CMB) taking into account simultaneous variation of other physical constants, namely the electron mass $m_{e}$ and the proton mass $m_{p}$. In other words, we consider the variation of Yukawa coupling and the QCD scale $\\Lambda_{\\mathrm{QCD}}$ in addition to the electromagnetic coupling. We clarify which parameters can be determined from CMB temperature anisotropy in terms of singular value decomposition. Assuming a relation among variations of coupling constants governed by a single scalar field (the dilaton), the 95\\,\\% confidence level (C.L.) constraint on $\\Delta\\alpha/\\alpha$ is found to be $-8.28 \\times 10^{-3} < \\Delta\\alpha/\\alpha < 1.81 \\times 10^{-3}$, which is tighter than the one obtained by considering only the change of $\\alpha$ and $m_{e}$. We also obtain the constraint on the time variation of the proton-to-electron mass ratio $\\mu \\equiv m_{p}/m_{e}$ to be $-0.52 < \\Delta\\mu/\\mu < 0.17$ (95\\,\\% C.L.) under the same assumption. Finally, we also implement a forecast for constraints from the PLANCK survey. ",
        "introduction": "Whether or not the physical constants are truly constant is not only a fundamental issue in physics but also an important probe of the theories of extra dimensions. Since the early studies of Dirac \\cite{Dirac:1937ti,Dirac:1938mt}, many theoretical models that could accomodate the variation of physical constants have been proposed \\cite{Bekenstein:1982eu,Barrow:2001iw,Barrow:2005qf,Uzan:2002vq, Chiba:2006xx}. Unification theories such as superstring theories predict the existence of an additional scalar field $\\Phi$ called ``dilaton\" and the destabilization of this field \\cite{Dine:1985he,Brustein:1992nk} might cause the variation of coupling constants. In more phenomenological contexts, dynamical scalar fields introduced to explain the recent discovery of the accerelated expansion may be non-minimally coupled to gauge fields in the Standard model, which may also lead to time varying coupling constants \\cite{Avelino:2006gc,Avelino:2008dc,Fujii:2007zg,Fujii:2008zj,Dent:2008vd,Avelino:2009fd}. Thus, investigating the variation of the coupling constants can be an important probe of these theories.  Not only these theoretical sides but also the observational analysis of various systems have been playing an important role in the investigation of the time variation of physical constants. Among them, the claim for the deviation of the past value of the fine structure constant $\\alpha$ from the present one by using the high-redshift quasar absorption system \\cite{Webb:1998cq,Webb:2000mn} had the great impact on this topic. Further analysis for the quasar absorption spectra gives the support for this positive result of a cosmological variation of $\\alpha$ \\cite{Murphy:2004,Murphy:2006vs,Murphy:2007qs,Molaro:2007kp}, though some authors takes objection to such variations \\cite{Chand:2004ct,Srianand:2004mq,Molaro:2007kp}.  Stimulated by these studies from both theoretical and observational sides, various constraints on the time variation of the fine structure constant $\\alpha$ have been investigated from diverse observations. We briefly summarize those terrestrial and celestial limits on $\\alpha$ as follows (see \\cite{Uzan:2002vq,Chiba:2001ui} for a review). The atomic clocks constrain the current value of the temporal derivative of $\\alpha$ as $\\dot{\\alpha}/\\alpha = (-3.3\\pm 3.0)\\times 10^{-16} \\ \\mathrm{yr}^{-1}$\\cite{Blatt:2008su,Fortier:2007jf,Peik:2004qn,Fischer:2004jt}. The measurement of the frequency ratio of aluminium and mercury single-ion optical clocks  provides $\\dot{\\alpha}/\\alpha=(-1.6\\pm 2.3) \\times 10^{-17} \\ \\mathrm{yr}^{-1}$ \\cite{Rosenband:2008}. From the analysis of Sm isotopes in the Oklo natural reactor in Gabon, we get two bounds on the variation of $\\alpha$ as $\\Delta\\alpha/\\alpha = -(0.8\\pm 1.0) \\times 10^{-8}$ and $\\Delta\\alpha/\\alpha=(0.88 \\pm 0.07) \\times 10^{-7}$ \\cite{Fujii:2002hc}, which measures the value at the redshift $z\\sim 0.1$ constraint. From the spectra of quasars, several limits have been obtained using various data sets: $\\Delta\\alpha/\\alpha = (-0.57\\pm 0.11) \\times 10^{-5} \\ (z \\sim 0.2-4.2$) from the Keck/HIRES instrument \\cite{Murphy:2004}, $\\Delta\\alpha/\\alpha = (-0.64 \\pm 0.36)\\times 10^{-5} \\ (z\\sim 0.4-2.3)$ from the Ultraviolet and Visual Echelle Spectrograph (UVES) instrument \\cite{Murphy:2006vs,Murphy:2007qs}. Big bang nucleosynthesis (BBN) provides constraints at very high redshifts ($z\\sim 10^{9}-10^{10}$), for example, $-5.0\\times 10^{-2} < \\Delta\\alpha/\\alpha < 1.0 \\times 10^{-2} \\ (95 \\% \\mathrm{C.L.})$ \\cite{Ichikawa:2002bt}. Finally, Cosmic Microwave Background (CMB) measures $\\alpha$ at $z\\sim 10^{3}$ and the limits from WMAP 1-year, 3-year and 5-year data read $-0.06 < \\Delta\\alpha/\\alpha < 0.01$ \\cite{Rocha:2003gc}, $-0.039 < \\Delta\\alpha/\\alpha < 0.010$ \\cite{Stefanescu:2007aa} and $-0.028 < \\Delta\\alpha/\\alpha < 0.026$ \\cite{Nakashima:2008cb} respectively, all of which are at $95\\% \\mathrm{C.L.}$ Recently, Menegoni et al. \\cite{Menegoni:2009rg} obtained more improved result $-0.013 < \\Delta\\alpha/\\alpha < 0.015$ at $95\\% \\mathrm{C.L.}$ from WMAP 5-year date combined with ACBAR, QUAD and BICEP experiments data.  Along with $\\alpha$, several observational constraints have been obtained on another physical constant, the proton-to-electron mass ratio $\\mu\\equiv m_{p}/m_{e}$. As in $\\alpha$, atomic clocks very tightly constrain the current value of its temporal derivative as $\\dot{\\mu}/\\mu = (-1.5\\pm 1.7)\\times 10^{-15} \\ \\mathrm{yr}^{-1}$\\cite{Blatt:2008su}. Reinhold et al. \\cite{Reinhold:2006zn} reports the non-vanishing variation of $\\mu$ as $\\Delta\\mu/\\mu = (2.4\\pm 0.6)\\times 10^{-5}$ from a weighted sum of accurate $\\mathrm{H}_{2}$ spectral lines in two quasar systems (corresponding redshifts are $z=2.59$ and $z=3.02$). Based on the quasar absorption spectra of $\\mathrm{NH}_{3}$, the limit is $\\Delta\\mu/\\mu = (0.6\\pm 1.9)\\times 10^{-6}$ at $z=0.6847$ \\cite{Flambaum:2007fa} and further detailed measurement by the same method tighten the bound as $\\Delta\\mu/\\mu = (0.74 \\pm 0.47)\\times 10^{-6}$ \\cite{Murphy:2008yy}.  In this paper, using WMAP 5-year data, we focus on the CMB constraint on $\\alpha$ in the case that the multiple physical constants may vary with time simultaneously according to the expectation values of a single dilaton field $\\Phi$. Although similar analyses has been done for BBN \\cite{Ichikawa:2002bt,Campbell:1994bf,Langacker:2001td,Dent:2001ga,Muller:2004gu,Dent:2008gx}, quasar absorption system \\cite{Calmet:2001nu,Calmet:2002ja} or the Oklo reactor \\cite{Olive:2002tz}, few attempts from CMB data have been made except for several preceding literatures \\cite{Ichikawa:2006nm,Landau:2008re,Scoccola:2008jw}, which treated simultaneous variations of only $\\alpha$ and $m_{e}$. In addition to $\\alpha$ and $m_{e}$, many theoretical models predict the time variation of the QCD scale $\\Lambda_{\\mathrm{QCD}}$ (which results from the variation of the strong coupling constant), so it is important to include the $\\Lambda_{\\mathrm{QCD}}$ effect in the analysis. We consider the effect of varying $\\Lambda_{\\mathrm{QCD}}$ as the change in the proton mass $m_{p}$ which is assumed to be proportional to $\\Lambda_{\\mathrm{QCD}}$ \\footnote{When we consider the time variation of $\\Lambda_{\\mathrm{QCD}}$, the neutron mass also changes. In this paper, we neglect the difference between the proton mass and the neutron mass, and when we refer to the time variation of the proton mass $m_{p}$, we implicitly include the time variation of the neutron mass assuming that the neutron mass changes in exactly the same way as the proton mass.} . Thus, we consider simultaneous variation of $\\alpha$, $m_{e}$ and $m_{p}$. Since how the variation of $\\Lambda_{\\mathrm{QCD}}$ is related to that of $\\alpha$ or $m_{e}$ depends on the details of the unification model \\cite{Dine:2002ir}, we here adopt a specific example of the relation among those physical constants following \\cite{Ichikawa:2002bt} or \\cite{Campbell:1994bf}, which is based on a low energy effective theory of string theory. The newly incorporated change of $\\Lambda_{\\mathrm{QCD}}$ leads to a new and tighter constraint on the time variation of $\\alpha$. We also obtain a constraint on $\\mu$ induced by the time dependence of $\\Lambda_{QCD}$ and $m_{e}$.  This paper is organized as follows. Section \\ref{sec:effect} provides how the variation of each physical constant affects the CMB power spectrum both qualitatively and quantitatively with particular attention to $\\Lambda_{\\mathrm{QCD}}$. Section \\ref{sec:dilaton} briefly describes the model we adopt based on string theory which governs the time evolution of the three physical constants we consider, namely, the variation of the fine structure constant, the electron mass and the proton mass. In section \\ref{sec:constraint}, we provide the limits of those variations. Section \\ref{sec:conclusion} is devoted to conclusion. The details of the calculations of the degeneracy and parameter dependence are given in the appendix A. In appendix B, we present a forecast for constraints from the PLANCK experiment. ",
        "conclusions": "\\label{sec:conclusion}  In summary, we have pointed out that the time variation of $\\Lambda_{\\mathrm{QCD}}$ can affect the CMB through the variation in the proton mass. We have shown that there is a notable effect when we impose the widely adopted assumptions of unification models found in the literature \\cite{Campbell:1994bf}. From the WMAP 5-year data, by MCMC analysis using the CosmoMC code, we have obtained the new and stringent constraint on the time variation of the fine structure constant as $-8.28\\times 10^{-3}<\\Delta\\alpha/\\alpha<1.81\\times 10^{-3}$. We have also verified that, in getting the above result, the newly considered effect of $\\Delta m_{p}/m_{p}$ is very important in that the degeneracy between $\\Delta m_{p}/m_{p}$ and $\\Omega_{B}$ is tremendously strong and that the small variation of $\\Delta\\alpha/\\alpha$ leads to the large variation of $\\Delta m_{p}/m_{p}$. Furthermore, by including the varying $m_{p}$, we have obtained a limit on the variation of proton-electron mass ratio $\\mu$ as $-0.52 < \\Delta\\mu/\\mu < 0.17$.  CMB observation will be much more precise to give a lot of information around the damping tail of the power spectrum in near future. The difference in the CMB power spectra between $\\Delta m_{p}/m_{p}$ and $\\Omega_{B}$ appears in the small scale diffusion damping because the damping scale is determined by the baryon number density $n_{B}$, probing the time variation of coupling constants including $\\Lambda_{\\mathrm{QCD}}$ by CMB will be even more promising. This is discussed more quantitatively in appendix A, and supported by a forecast for future CMB survey in appendix B.   \\bigskip"
    },
    "0910/0910.5484_arXiv.txt": {
        "abstract": "Radio, infrared, and optical observations of the 2006 eruption of the symbiotic recurrent nova RS Ophiuchi (RS Oph) showed that the explosion produced non-spherical ejecta. Some of this ejected material was in the form of bipolar jets to the east and west of the central source.  Here we describe X-ray observations taken with the $Chandra$ X-ray Observatory one and a half years after the beginning of the outburst that reveal narrow, extended structure with a position angle of approximately 300 degrees (east of north). Although the orientation of the extended feature in the X-ray image is consistent with the readout direction of the CCD detector, extensive testing suggests that the feature is not an artifact. Assuming it is not an instrumental effect, the extended X-ray structure shows hot plasma stretching more than 1,900 AU from the central binary (taking a distance of 1.6 kpc). The X-ray emission is elongated in the northwest direction --- in line with the extended infrared emission and some minor features in the published radio image. It is less consistent with the orientation of the radio jets and the main bipolar optical structure. Most of the photons in the extended X-ray structure have energies of less than 0.8 keV. If the extended X-ray feature was produced when the nova explosion occurred, then its 1$^{\\prime\\prime}$.2 length as of 2007 August implies that it expanded at an average rate of more than 2 mas d$^{-1}$, which corresponds to a flow speed of greater than 6,000 km/s (d/1.6 kpc) in the plane of the sky. This expansion rate is similar to the earliest measured expansion rates for the radio jets. ",
        "introduction": "Novae are the most common major stellar explosions in the universe. The basic properties of these explosions are well explained by models in which a thermonuclear runaway at the base of the envelope of an accreting white dwarf (WD) causes the WD to brighten to near the Eddington luminosity and eject its envelope in a spherical outflow. An increasing number of observations, however, are indicating that the ejecta from many novae are anything but spherical \\citep[e.g.,][]{slavin95,gill00,krautter,harman03}, even in the X-rays \\citep[detected a century after outburst by][]{balman}. The properties of the outflows triggered by nova explosions on massive WDs are particularly important because they may influence or reveal the amount of mass lost per explosion, and thereby the ability of the stars to become type Ia supernovae.   Since the recurrence time for a nova explosion in a given system is inversely related to the mass of the WD \\citep{yaron}, interacting binaries that contain WDs with masses greater than approximately 1.1 M$_{\\odot}$ can have recurrence times of less than 100 years; these systems are the so-called recurrent novae. In symbiotic recurrent novae -- of which only 4 are known: \\rsoph, T CrB, V3890 Sgr and V745 Sco -- the accreting WD is embedded in a dense nebula fed by the wind of the red-giant companion. RS Oph went into outburst for the sixth recorded time (previous outbursts occurred in 1898, 1933, 1958, 1967, 1985) on February 12th 2006 \\citep{narumi06}. The astronomical community observed this outburst at virtually every wavelength; the X-ray range was extremely well covered, with observations by $RXTE$ \\citep{sokoloski06}, $Swift$ \\citep{hachisu07,bode06}, $Chandra$, and $XMM$-Newton \\citep{ness07,nelson08}.  Observations of \\rsoph indicate that the morphology of the expanding shell is complex and highly asymmetric. Within a few weeks of the start of the 2006 eruption, X-ray spectra showed that the blast-wave evolution deviated from self-similar, spherical expansion \\citep{sokoloski06,bode06}. Bipolar structure along the east-west direction was clearly seen in radio images taken a few weeks into the 2006 outburst \\citep{obrien06,rupen08}; several months after the start of the outburst, observations showed that the radio jets were composed of highly collimated outflows that powered synchrotron-emitting lobes \\citep{jeno08}. At optical wavelengths, % [Ne V]$\\lambda$3426 images 5 months after the start of the outburst revealed an expanding nebular remnant with a double-ring structure \\citep{bode07}. The major axis of the double ring was also oriented east-west.  Near infrared (NIR) interferometric observations also revealed extended asymmetric structure in the early expanding shell --- but along a different direction. On day 5.5 after the start of the outburst, \\citet{chesneau07} found that the NIR emission arose from a rapidly expanding elliptical region with major axis oriented from northwest to southeast. NIR observations in the first month of the outburst by \\citet{lane07} also revealed an asymmetric structure with a position angle (P.A.) of approximately 120$^{\\circ}$/300$^{\\circ}$.    In this letter, we describe the first detection of extended X-ray structure from RS Oph. Taken one and half years after the start of the 2006 eruption, our $Chandra$ observations reveal a narrow X-ray feature along the direction of the NIR extended emission. We describe the observation in section \\S\\ref{sec:obs}, detail the analysis and results in section \\S\\ref{sec:image}, and discuss possible interpretations of the results in section \\S\\ref{sec:discuss}. ",
        "conclusions": "\\label{sec:discuss}  Assuming, as suggested by our extensive testing, that the extended X-ray structure is not an artifact, the properties of this feature indicate that it consists of a few tenths of a percent to a few percent of the ejecta mass, and that the outermost X-ray emitting material is moving at a speed approximately half of the escape velocity from the WD (the escape velocity is approximately 12,600 km s$^{-1}$ for a 1.35 M$_{\\odot}$ WD). If the material comprising the narrow X-ray feature was ejected at the time of the nova explosion, the length of the X-ray structure suggests an average expansion rate in the plane of the sky of $\\sim$2.3 mas d$^{-1}$. This rate is similar to the early expansion rate of $\\sim$ 2 mas d$^{-1}$ for the radio jet \\citep{jeno08}. The rate of expansion of the X-ray feature corresponds to velocities in the plane of the sky of about $\\sim6,300\\pm 1600$~km $s^{-1}$ ($d$/1.6 kpc). Since the X-ray emitting material is unlikely to be moving exactly in the plane of the sky, the true velocity could be higher.  A strong shock moving at such a high speed would be expected to produce hard X-ray emission (with $kT \\simeq$ 40 keV in the strong-shock case); however we find only soft X-ray emission from the extended feature (\\S\\ref{sec:struc}).  The discrepancy between the expected temperature in the strong shock case and the observed photon energy is also present in the diffuse X-ray emission from planetary nebulae \\citep{soker03} and the two symbiotic stars with X-ray jets -- CH~Cyg \\citep{galloway04} and R~Aqr \\citep{kellog07}.  The narrow X-ray feature is most extended in the 0.3--0.8 keV energy band and contains approximately 40 photons in this band. A crude estimation of its mass can be made if we assume that the emission is due to an absorbed, optically thin thermal plasma with $kT$=0.5 keV and n$_{H}$=1.8$\\times$10$^{21}$ cm$^{-2}$ \\citep{bode06}. We use PIMMS to obtain the observed emission measure of such a plasma and, assuming that the narrow X-ray feature has a radius of $\\sim$ 0$^{\\prime\\prime}$.9 and length of 1$^{\\prime\\prime}$.2 (Fig.~\\ref{fig3}), find that the density is $\\sim$40 cm$^{-3}$ and the mass of the extended X-ray emitting material is $\\sim$10$^{-8}$ M$_{\\odot}$ (assuming that there is one proton for every free electron). For comparison, estimates of the amount of material ejected immediately after the outburst range from $\\sim$10$^{-7}$ M$_{\\odot}$ to $\\sim$10$^{-6}$ M$_{\\odot}$ \\citep{sokoloski06,hachisu07}.  The orientation of the narrow X-ray feature is less consistent with that of the radio jets \\citep{obrien06,rupen08,jeno08} than with other features detected at radio and infrared (NIR) wavelengths. Early radio observations showed that the ring of synchrotron-emitting plasma associated with the expanding blast wave had a bright spot at a P.A.$\\approx$120$^{\\circ}$ \\citep{obrien06,rupen08}. If one were to superpose the X-ray and radio images and draw a line along the extended X-ray structure, it would intersect the bright spots on the radio ring. Moreover, the radio ring was actually more of a C-shape, with the opening in the ring at the same P.A. as the X-ray feature. NIR interferometry uncovered extended emission with a P.A. of 120$^{\\circ}$/300$^{\\circ}$ from day 5.5 to a few months after the nova explosions \\citep{chesneau07, lane07}, and \\citet{chesneau07} suggested a possible connection between the radio and IR morphologies. Our detection of an X-ray feature extending in the northwest direction, like features in the radio and NIR, suggests that the structures along this direction could all be related (see Fig.~\\ref{fig4}).  Although the X-ray structure has some similarities to features predicted by models of the explosion, many properties of the observed narrow X-ray feature are unexplained, and the true origin of this extended X-ray emitting gas remains unclear.  Computational modeling suggests that the blast-wave expansion should be strongly asymmetric. For example, in the 3-dimensional simulations of \\citet{walder08}, the blastwave moves into a highly inhomogeneous pre-outburst density distribution and forms an elongated, bipolar structure perpendicular to the plane of the orbit.  We do not see any extended X-ray emission in this direction. In contrast, models by \\citet{orlando09} predict a one-sided X-ray structure -- as we have found -- due to the reflection of the blastwave off of the red giant. Unfortunately, however, the position of the red giant at the time of the nova explosion \\citep[according to the ephemeris of][]{brandi09} would have led to an X-ray extension in the north-south direction, which is inconsistent with our detection of extended X-ray emission at a P.A. of ~300 degrees. Thus, our results suggest that either the orbital solution for this system needs to be re-examined or that the explosion somehow ejected X-ray emitting plasma in a direction that was neither perpendicular to the plane of the orbit nor along the line joining the WD and the red giant.   \\begin{figure*} \\begin{center} \\includegraphics[scale=0.313]{f4a.eps} \\includegraphics[height=6.3cm]{f4b.ps} \\caption{Comparison of RS~Oph images in radio, infrared and X-rays. {\\it Left:} an image of radio synchrotron-emitting plasma with NIR contours overlaid \\citep{chesneau07}. {\\it Right:} the X-ray image obtained after SER and PSF-deconvolution processing (same as panel (d) in Fig.~\\ref{fig1}). The X-ray feature extends along the same direction as the NIR and radio features, suggesting that they could all be related. Note that the two panels show very different scales.} \\label{fig4} \\end{center} \\end{figure*}"
    },
    "0910/0910.0432_arXiv.txt": {
        "abstract": "The photometric and spectroscopic observational campaign organized for the 2008/9 eclipse of EE~Cep revealed features, which indicate that the eclipsing disk in the EE~Cep system has a multi-ring structure.  We suggest that the gaps in the disk can be related to the possible planet formation. ",
        "introduction": "The eclipses of the 11th magnitude star EE Cep have been observed with a~period of 5.6 yr from the early 1950-ies.  Their depths change in a wide range from about $0\\fm5$ to $2\\fm0$ \\citep[see][]{Gra2003}, however all of them show the same features: they are almost gray and have the same asymmetric shape (descending branch of every eclipse has longer duration than the ascending one).  In all eclipses it is possible to distinguish five phases (shown in Fig.~1): ingress (1--2) and egress (3--4), respectively preceded and followed by the extended atmospheric parts (1$_{\\rm a}$--1 and 4--4$_{\\rm a}$), and the slope-bottom transit (2--3).  The most plausible explanation of the observed photometric behavior was proposed by \\citet{MiGr1999}.  They suggested that the secondary consists of a dark disk, with opaque interior and semi-transparent exterior, around a low luminosity central object.  The inclination of the disk to the line of sight and the tilt of its cross-section to the direction of motion are changed by precession.  This causes the changes in the depth and the duration of the eclipses.  A significant impact parameter is responsible for the observed asymmetry of the eclipses.  This model can explain the shallow ($0\\fm6$), flat-bottomed eclipse observed in 1969 if we assume a nearly edge-on and non-tilted projection~of~the~disk.  \\begin{figure}[!ht] \\begin{center} \\includegraphics[angle=0, width=0.64 \\textwidth]{Galan1fig1.eps} \\end{center} \\caption{\\small A schematic illustration of the eclipses in the EE~Cep system.  The contact times ($1_{\\rm a},1,2,3,4,4_{\\rm a}$) distinguished in the light curve (right) refer to characteristic configuration of the disk and the star during~the~eclipse~(left).} \\end{figure} ",
        "conclusions": "A comparison of the photometric and spectroscopic data obtained during the last two eclipses (see Fig.~6) reveals some new characteristics of the eclipsing disk in the EE~Cep system.  The durations of last two eclipses were longer than we expected (about 90 days), and they both were preceeded and followed by very shallow minima which are perhaps repeatable in each orbital cycle.  The 2008/9 campaign results confirmed the existence of a gap in the disk.  Additionaly, the data present some indications of the existence of a~second, outer gap in the disk.  The possible multi-ring structure of the eclipsing disk in EE~Cep suggests the existence of some massive bodies that could be responsible for their formation.  This means that we may observe signs of planet formation in a circumstellar disk in EE~Cep.  The results presented here show that the disk in EE~Cep system is very similar to the multi-ring structure observed in $\\varepsilon$ Aurigae \\citep{Fer1990}.  During the EE~Cep eclipses additional components appear in the $H_{\\alpha}$ emission line and in the NaI absorption doublet.  Towards the mid-eclipse an absorption component appears and grows in the $H_{\\alpha}$ profile and during the minimum it is very deep and broad.  The sodium doublet line profiles evolve during the eclipse and in the minimum multi-component structure with at least two additional absorption components can be seen. During the shallow minimum at orbital phase 0.017 an additional NaI absorption component was also present, while it was absent soon after the egress.  The spectra from the last two eclipses suggest that the absorption lines evolution is the same during each cycle.  Unfortunately spectroscopic observations coverage of the last eclipse was very bad.  Many more spectroscopic observations during the next eclipse would be needed to understand the nature of EE~Cep.  Photometry in the infrared {\\it JHK} passbands and the radial velocity variations of the hot component could~be~invaluable~as~well.  \\begin{figure}[!ht] \\begin{center} \\includegraphics[angle=0, width=0.98 \\textwidth]{Galan1fig6.eps} \\end{center} \\caption{\\small Photometry and spectra from last two eclipses are compared. {\\textit {Left:}} The observations between orbital phase 0.9 and 0.1 are presented.  {\\textit {Right:}} An expanded view of 0.04 of orbital phase marked with vertical dashed lines in the left panel is shown.  {\\textit {Top:}} Photometry of the 2008/9 eclipse.  {\\textit {Bottom:}} Photometry of the 2003 eclipse.  The times close to the very shallow minima which occur in both eclipses at phase near of $\\pm0\\fp017$ are denoted with arrows. {\\textit {Middle:}} The evolution of the NaI absorption and the H$_{\\alpha}$ emission lines.  The spectra obtained during and close to the 2003 eclipse are given with dashed lines, while those obtained during the last eclipse, with solid lines.  The positions of the continuum levels refer to the orbital phases.} \\end{figure}"
    },
    "0910/0910.1602_arXiv.txt": {
        "abstract": "Dark matter charged under a new gauge sector, as motivated by recent data, suggests a rich GeV-scale ``dark sector'' weakly coupled to the Standard Model by gauge kinetic mixing.  The new gauge bosons can decay to Standard Model leptons, but this mode is suppressed if decays into lighter ``dark sector'' particles are kinematically allowed. These particles in turn typically have macroscopic decay lifetimes that are constrained by two classes of experiments, which we discuss. Lifetimes of $10 \\cm \\lesssim c\\tau \\lesssim 10^{8} \\cm$ are constrained by existing terrestrial beam-dump experiments.  If, in addition, dark matter captured in the Sun (or Earth) annihilates into these particles, lifetimes up to $\\sim 10^{15} \\cm$ are constrained by solar observations.  These bounds span fourteen orders of magnitude in lifetime, but they are not exhaustive.  Accordingly, we identify promising new directions for experiments including searches for displaced di-muons in B-factories, studies at high-energy and -intensity proton beam dumps, precision gamma-ray and electronic measurements of the Sun, and milli-charge searches re-analyzed in this new context. ",
        "introduction": "Recent astrophysical anomalies have motivated models \\cite{Finkbeiner:2007kk,Pospelov:2007mp,Cholis:2008vb,ArkaniHamed:2008qn,Pospelov:2008jd} of TeV-scale dark matter interacting with new states at the GeV scale. A compelling candidate is a new gauge force, mediated by a $\\GeV$-mass vector-boson that mixes kinetically with hypercharge: \\bea \\label{eq:Lag} \\mathcal{L} &=&  \\mathcal{L}_{\\textrm{{SM}}}  + i\\bar{\\chi}\\left(\\slashed{\\partial} + i g \\slashed{A'}\\right) \\chi + M_\\chi\\bar{\\chi}\\chi \\\\\\nonumber &-&\\frac{1}{4}F'^{,\\mu\\nu}F'_{\\mu\\nu}+ \\frac{\\epsilon}{2 \\cos\\theta_W} F'_{\\mu\\nu} B^{\\mu\\nu} + \\mathcal{L}_{\\textrm{{Dark}}}. \\eea A program of searches in high-luminosity $\\epm$ collider data \\cite{Essig:2009nc,Batell:2009yf,Reece:2009un}, new fixed-target experiments \\cite{Bjorken:2009mm,Batell:2009di}, and high-energy colliders \\cite{ArkaniHamed:2008qp,Baumgart:2009tn} can search quite exhaustively for the new gauge boson $A'$ if it decays into a lepton pair through the $\\epsilon$-suppressed kinetic mixing term.  The $A'$ can, however, be accompanied by ``dark sector'' particles (other vectors, Higgs-like scalars, pseudoscalars, or fermions) to which it couples without $\\epsilon$-suppression, so that the decays of $A'$ into dark sector particles dominate. These particles may decay to Standard Model matter (and indeed some \\emph{must}, if dark matter annihilation is to explain the astrophysical anomalies), but unlike the $A'$ their lifetimes are typically macroscopic --- we will call such particles ``long-lived particles'' (LLP's).  \\begin{figure} \\includegraphics[width=0.45\\textwidth]{figures/BasicSetup.pdf} \\caption{A pictorial summary of the two similar approaches discussed in this paper: A strong source, either a beam dump or dark matter annihilation in the Sun, produces dark sector particles. LLP's can penetrate the shield, either the Sun's interior or the beam dump shielding. Decays downstream over a baseline, $L_{\\mbox{Baseline}}$, can be detected.} \\label{fig:Setup} \\end{figure}  In this paper, we discuss observational constraints on the decays of the LLPs which cover lifetimes spanning fourteen orders of magnitude, $10 \\cm \\lesssim c\\tau \\lesssim 10^{15} \\cm$. A schematic diagram illustrating the type of experiments we consider is shown in Fig. \\ref{fig:Setup}. The constraints fall into two categories: \\begin{description} \\item[Terrestrial beam dump experiments] (Sec.~\\ref{sec:beamDump}): Experiments at proton and electron beam dumps search for decays of long-lived particles in an instrumented region downstream of a thick shield.  Though $A'$ production scales as  $\\epsilon^2$, the integrated luminosities $\\sim10^2\\ab^{-1}$ allow searches for LLP's with mass $\\lsim\\rm{few }\\GeV$ and $10 \\lesssim c\\tau \\lesssim 10^8$ cm. \\item[Annihilation of Dark Matter in the Sun] (Sec. ~\\ref{sec:solar}): Dark matter captured in the Sun through $A'$-mediated scattering annihilates efficiently, with a rate scaling as $\\epsilon^2$. Annihilations into LLP's lighter than the $\\tau$ threshold result in $e$, $\\mu$, and light hadrons, which can be detected when the LLP decays outside the Sun but before passing the Earth.  In this case, neutrino, electron, and gamma ray observations together constrain dark matter annihilation into LLP's with $10^7 \\cm \\lesssim c\\tau \\lesssim 10^{15}$ cm. \\end{description} In Sections \\ref{sec:beamDump} and \\ref{sec:solar}, we describe the best available limits.  In Section \\ref{sec:Example}, we interpret these limits in a particular illustrative case: a Higgs-like scalar in the dark sector that decays into two Standard Model fermions.  We conclude in Section \\ref{sec:discussion}, and highlight the value of direct $A'\\rightarrow \\ell^+\\ell^-$ searches, searches for displaced di-muons in B-factories, high-energy and high-intensity proton beam dumps, precision gamma-ray and electronic measurements of the Sun, and milli-charge searches in the context of these models.  \\subsection{Decay Lengths in Dark Sectors}\\label{sec:orientation} Before proceeding, we summarize the parametric scaling that governs the decays of various particles in a dark sector which is coupled to the Standard Model through gauge kinetic mixing.  In non-Abelian dark sectors, multiple gauge bosons can mix with one another \\cite{ArkaniHamed:2008qn,Baumgart:2009tn}, so that each acquires an effective kinetic mixing $\\epsilon_{eff}$ and lifetime $\\propto \\epsilon_{eff}^{-2}$.  Particles $\\chi$ of other spins can have three-body decays into Standard Model charged matter ($\\chi \\rightarrow f^+ f^- \\chi'$ through an off-shell $A'$), with lifetimes $\\propto \\epsilon^{-2}$ but longer than those of vectors due to three-body phase space.  For $\\epsilon \\sim 10^{-2}-\\mbox{ few}\\times10^{-5}$ and moderate mass hierarchies these lifetimes are less than 1 m, and multi-body final states can be observed in a collider detector. B-factory searches in higher-multiplicity final states (e.g. \\cite{Aubert:2009pw}) can be quite sensitive to these decay modes \\cite{Batell:2009yf,Essig:2009nc}.  Much longer lifetimes can arise from $\\epsilon^{-4}$ scaling or decays through higher-dimension operators --- these are the modes of greatest interest for our discussion.  In particular, we consider a Higgs-like boson $h_D$ below the $A'$ mass.  The $A'$ can decay into $h_D$ only in models with multiple dark-sector Higgses, as the second decay product must be a pseudoscalar.  Several different parametric decay lifetimes are possible.  The dark Higgs can decay through gauge kinetic mixing, with $\\epsilon^{-4}$ scaling, or faster if it mixes with Standard Model Higgs bosons (e.g. through a quartic term $\\epsilon_\\lambda |h_D|^2 |h^2|$).  If the Dark Higgs is stabilized by an accidental symmetry so that its decay is mediated by higher-dimension operators as noted in \\cite{Katz:2009qq,Morrissey:2009ur}, the lifetimes can be significantly longer.  For suitable parameter choices, all three decay scenarios (kinetic mixing, higgs mixing, and higher-dimension operators) lead to lifetimes $10 \\cm \\lesssim c\\tau \\lesssim 10^{15}$ cm, so that beam-dump and solar limits apply.  The kinetic mixing decay is a simple and illustrative benchmark, because the decay lifetime and terrestrial production cross-sections both depend on $\\epsilon$ and the particle masses, with no additional free parameters.  After discussing the model-independent limits, we will consider this case explicitly in Section \\ref{sec:Example}. ",
        "conclusions": "\\label{sec:discussion} Dark sectors coupled to the Standard Model through gauge kinetic mixing may contain an array of macroscopically long-lived particles (LLP's) in addition to the kinetically mixed gauge boson.  Gauge bosons can escape detection in accelerator experiments if they decay into LLP's with lifetimes between 10 cm and $10^{15}$ cm, rather than to Standard Model fermions.  LLPs with lifetimes in this range would escape collider detectors unobserved, but are sufficiently short-lived that BBN constraints are largely irrelevant.  Remarkably, two classes of data constrain these scenarios across these fourteen decades in lifetime.  The proton beam dump experiment CHARM places strong constraints on scalar LLP's with $10 \\lesssim c\\tau \\lesssim 10^8$ cm decaying into muons or electrons. If dark matter annihilates into dark sector LLP's, then Solar capture of dark matter results in limits on LLP's with $10^7 \\cm \\lesssim c\\tau \\lesssim 10^{15}$ cm for similar decay modes.  These limits each depend on different assumptions and production mechanisms, so they cannot be combined into a robust exclusion. Nonetheless, the case of a Higgs-like boson in the dark sector decaying through kinetic mixing allows an illustrative comparison.  The sensitivity of solar limits --- when applicable --- nearly matches on to the parameter region excluded by CHARM. Together, these constraints disfavor $A'$ decays to scalars, pseudoscalars and other long-lived particles in the parameter region best suited to explain the DAMA modulation signal through iDM \\cite{Katz:2009qq,Alves:2009nf,Morrissey:2009ur}.  In particular, if LLP's decay through higher-dimension operators, lowering the mass scale of these operators to evade BBN constraints may bring them into conflict with experimental data.  Given these limits and the remaining unconstrained parameter regions, our analysis suggests several promising searches with existing data and small-scale experiments: \\begin{description} \\item[Prompt Vector Decays to $\\ell^+\\ell^-$:] Remarkably, for $\\epsilon \\sim 10^{-5}-10^{-3}$, $A'$ decays into LLPs are \\emph{much more} constrained than decays directly into Standard Model states.  This observation compounds the motivation for a sharply defined search program to constrain direct $A'$ decays to $\\ell^+\\ell^-$ \\cite{Bjorken:2009mm,Reece:2009un}. \\item[Searches in $e^+e^-$ Collider Data:] LLP lifetimes too short to be observed in a beam-dump experiment ($\\gamma c\\tau \\lesssim 10$m at $E_{LLP}\\sim 50 \\GeV$, e.g. region I of Fig.~\\ref{fig:higgsConstraints}) give rise to observable displaced decays within the detectors of  $e^+e^-$ collider experiments such as Babar, Belle, KLOE, and BES-III. Moreover, such lifetimes arise from $\\epsilon \\gtrsim 10^{-3}$ for which many thousands of events are possible in $\\gamma + A'$ and $h_D A'$ production.  The most visible such decay may be into $\\mu^+\\mu^-$ (or $\\pi^+\\pi^-$) pairs reconstructing a displaced vertex, which may be accompanied by a photon, invisible particles, and/or additional $e$, $\\mu$, or $\\pi$ pairs. \\item[High-Energy and High Intensity Beam Dumps:] High-energy beam dumps with cumulative charges on target exceeding a coulomb may extend sensitivity into region II of Fig. \\ref{fig:higgsConstraints}.  Neutrino factories with near detectors such as MINOS have geometries well-suited for beam-dump LLP searches \\cite{Batell:2009di,Essig:2009aa}, and integrated luminosity $\\sim 10^3$ larger than that of CHARM. \\item[Solar Gamma Ray and Electron Measurements:] LLP lifetimes $\\sim 10^{3}-10^{5}$ seconds (region III in Fig. \\ref{fig:higgsConstraints}) are too short to disrupt light element abundances, but only a fraction of LLP's can decay between the Sun and the Earth, so very sensitive solar measurements are needed to probe these scenarios.  Solar gamma-ray observations above 10 GeV from ACTs and FERMI can explore this region, as can searches for electronic signals correlated with the Sun \\cite{Schuster:2009aa}. \\item[Milli-Charged Particle Searches:] Electron beam-dump experiments searching for milli-charged particles (e.g. \\cite{Prinz:1998ua,Jaros:1995kz}) may be sensitive to LLP scattering off a nuclear target $N$ in the active detector volume, such as $h_D+N\\rightarrow A'+N$.  Searches in existing data may be sensitive to $\\epsilon \\gtrsim 10^{-3} (100 \\MeV/m_{A'})$ and $m_{LLP} \\lesssim 100$ MeV. \\end{description}   {\\it Note added: While this work was brought to completion two related works} \\cite{Batell:2009zp}, {\\it and } \\cite{Meade:2009mu}{\\it appeared.}"
    },
    "0910/0910.4161_arXiv.txt": {
        "abstract": "It is generally taken for granted that reionization has completed by $z=6$, due to the detection of flux in the \\lya\\ forest at redshifts $z<6$.  However, since reionization is expected to be highly inhomogeneous, much of the spectra pass through just the ionized component of the intergalactic medium (IGM) even for non-negligible values of the volume-weighted mean neutral hydrogen fraction, $\\avenf$.  We study the expected signature of an incomplete reionization at z$\\sim$ 5--6, using very large-scale (2 Gpc) seminumeric simulations.  We find that ruling out an incomplete reionization is difficult at these redshifts since: (1) quasars reside in biased regions of the ionization field, with fewer surrounding HI patches than implied by the global mean, $\\avenf$; this bias extends tens of comoving megaparsecs for $\\avenf\\lsim0.1$; (2) absorption from the residual neutral hydrogen inside the ionized IGM generally dominates over the absorption from the remaining HI regions, allowing them to effectively ``hide'' among the many dark spectral patches; (3) modeling the \\lya\\ forest and its redshift evolution even in just the ionized IGM is very difficult, and nearly impossible to do a priori.  We propose using the fraction of pixels which are dark as a simple, nearly model-independent upper limit on $\\avenf$.  Alternately, the size distribution of regions with no detectable flux (dark gaps) can be used to place a more model dependent constraint.  Either way, the current sample of quasars is statistically insufficient to constrain $\\avenf$ at $z\\sim6$ to even the 10 per cent level.  At $z\\sim5$, where there are more available sightlines and the forest is less dark, constraining $\\avenf\\lsim0.1$ might be possible given a large dynamic range from very deep spectra and/or the \\lyb\\ forest. We conclude with the caution against over-interpreting the observations.  There is currently no direct evidence that reionization was complete by $z\\sim$ 5 -- 6. ",
        "introduction": "\\label{sec:intro}  Theoretical arguments suggest that the first astrophysical objects ignited at redshifts $z \\lsim$ 30--40.  The ionizing radiation of these early stars and black holes carved out bubbles of ionized hydrogen around them.  As subsequent generations of objects formed, these HII regions gradually merged to create a complex HII morphology.  This epoch of reionization terminated in so-called ``overlap'', when the myriad ionization fronts merged, and the only remaining pockets of hydrogen with a substantial neutral fraction were dense, self-shielded systems, such as Lyman limit systems (LLSs) and damped Lyman alpha systems (DLAs).  Reionization is exceedingly complex, and involves uncertain, observationally-underconstrained astrophysics.  Early interpretations of observations generally made the simplifying assumption that reionization is spatially homogeneous.  However, an extremely inhomogenous reionization seems to be an unavoidable conclusion from recent analytic models \\citep{FZH04}, as well as large-scale numeric (e.g. \\citealt{CSW03, Sokasian03, Iliev06_sim, McQuinn07, TC07}), and seminumeric \\citep{Zahn07, MF07, GW08, Alvarez08, Thomas09} simulations.  The patchiness of reionization has many important consequences on the interpretation of observations, prompting calls for caution when interpreting \\lya\\ damping wing absorption in high-$z$ spectra \\citep{MF08damp, McQuinn08}, redshift evolution of the Lyman alpha emitter (LAE) luminosity function \\citep{FZH06, McQuinn07LAE, MF08LAE, Iliev08}, the apparent size of the proximity region in high-$z$ quasar spectra \\citep{Lidz07, AA07}, and a sharp redshift evolution in the effective Gunn-Peterson (GP; \\citealt{GP65}) optical depth, $\\taueff$, \\citep{FM09}.  Despite its complexity, a hard lower limit for the overlap phase of reionization of $z\\sim6$ is often quoted in the literature.  This result is justified by the detection of flux in the \\lya\\ GP throughs of $z\\lsim6$ quasars (QSOs).  Because of the high resonance optical depth of the \\lya\\ line, the detection of flux implies that the HI fraction in the intergalactic medium (IGM) is less than $\\sim10^{-4}$ at these redshifts  (e.g. \\citealt{Fan01, Becker01, Djorgovski01, SC02}).  Although this number might be appropriate for the residual HI in the {\\it ionized component} of the IGM, when viewed in the context of a patchy reionization the picture becomes less clear; in fact, the step from ``some flux is detected'' to ``overlap has occurred'' might require a large leap of (reionization model based) faith.  The interpretation of the QSO \\lya\\ forest in this regime is complicated for several reasons.  Firstly, as reionization nears completion, the remaining neutral patches become increasingly rare.  These neutral patches also have very large line-of-sight (LOS) fluctuations, which are larger in 1D than in the analogous 3D volume (c.f. \\citealt{LOF06} for a related treatment of the sightline-to-sightline fluctuations arising from just the density field).  Large sections of the line-of-sight therefore should pass through ionized IGM even for non-negligible values of the volume averaged hydrogen neutral fraction, $\\avenf\\lsim0.1$.  Further complicating matters is the fact that QSO host halos are extremely biased, and preferentially reside close to the centers of large HII regions, or large scale regions with fewer HI patches than implied by the global mean neutral fraction (e.g. \\citealt{Lidz07, AA07, GW08}).  Thus, one would have to look at spectral regions distant from the emission line center to get a representative sample of the ionization state of the Universe.  At these high redshifts ($z\\sim$ 5--6), the GP troughs are already very dark with large sightline-to-sightline fluctuations {\\it even from the residual neutral hydrogen in HII regions assuming a uniform UV background} (UVB; e.g. \\citealt{MHR00, Fan06, LOF06}).  Therefore, it is quite challenging to statistically distinguish in the \\lya\\ spectra the {\\it additional} dark patches arising from pre-overlap neutral regions, from the dark patches already present due to the residual neutral hydrogen, LLSs and DLAs in the ionized IGM.  Indeed, \\citet{Lidz07} point out that the final overlap stage of reionization might have occurred at $z<6$, given these difficulties in interpreting the observed spectra.  Here we quantify and further explore this scenario.  Specifically, {\\it the purpose of this paper is to quantify how confident can we be from current \\lya\\ forest data that overlap did not occur at $z<6$, given that reionization is very inhomogenous}.  For this purpose, we make use of the seminumeric simulation, DexM\\footnote{http://www.astro.princeton.edu/$\\sim$mesinger/DexM/}, which allows us to efficiently generate density, halo, and ionization fields on very large scales and with a large dynamic range.  This is an essential property of reionization studies, especially those involving rare, massive QSOs, since one must be able to statistically capture the ionization field, resolve $M\\gsim$ $10^{12}$--$10^{13}$ $\\Msun$ QSO host halos, while including ionizing photons from very small halos capable of efficient atomic hydrogen cooling ($\\Tvir \\gsim 10^4$ K, or $M \\gsim 10^8 \\Msun$ for the redshifts in question).  In \\S \\ref{sec:sims}, we briefly summarize the large-scale seminumeric simulations we use for this study.  In \\S \\ref{sec:bias}, we show how QSO host halos locations are biased with respect to the ionization field.  In \\S \\ref{sec:excess}, we present simple, general estimates of the excess absorption in the \\lya\\ forest from incomplete, patchy reionization.  In \\S \\ref{sec:dark_gaps}, we include the geometry of HI patches, and present distributions of spectral dark gaps.  Finally in \\S \\ref{sec:conc}, we summarize our conclusions.  We quote all quantities in comoving units.  We adopt the background cosmological parameters ($\\Omega_\\Lambda$, $\\Omega_{\\rm M}$, $\\Omega_b$, $n$, $\\sigma_8$, $H_0$) = (0.72, 0.28, 0.046, 0.96, 0.82, 70 km s$^{-1}$ Mpc$^{-1}$), matching the five--year results of the {\\it WMAP} satellite \\citep{Dunkley09, Komatsu09}.  Throughout the paper, we refer to overlap as the ``end'' of reionization.  As mentioned earlier, the end of reionization is somewhat difficult to define, since even after overlap a few percent of the {\\it mass weighted} HI fraction remains in dense absorption systems such as LLSs, DLAs and galaxies (e.g. \\citealt{PW09}).  These systems generally correspond to overdensities (e.g. \\citealt{MHR00}), unlike the HI patches pre-overlap, which preferentially reside in voids due to the ``inside-out'' nature of reionization on large scales.  Furthermore, these absorption systems are fairly small (e.g. \\citealt{Schaye01}) occupying a modest fraction of the total volume of the Universe. Here we use the {\\it volume weighted} neutral hydrogen fraction, $\\avenf$, as a proxy for the reionization status.  Thus we define ``post-overlap'' as the regime where $\\avenf \\rightarrow 0$, i.e. when the only remaining neutral regions are in dense, self-shielded systems. ",
        "conclusions": "\\label{sec:conc}  Much controversy has surrounded various claims concerning the ionization state of the Universe at $z>6$, as probed by the spectra of high redshift quasars. Ironically, the same arguments used to claim that the spectra of these quasars are consistent with reionization completing by $z>>6$ can also be used to claim the spectra are consistent with reionization {\\it not} completing by $z<6$.  In this work, we re-examine the often-quoted evidence of $\\avenf\\lsim10^{-4}$--$10^{-3}$ from the detection of flux in the spectra of $z\\lsim6$ QSOs.  Because reionization is very patchy, and the remaining neutral patches become very rare as reionization progresses, it is difficult to claim reionization was complete by this redshift \\citep{Lidz07}.  To quantify these difficulties, we take advantage of the large dynamic range provided by the seminumerical modeling tool, DexM.  Our simulation boxes are 2 Gpc on a side, resolve the likely host halos of QSOs, and generate ionization fields accounting for ionizing photons from halos with virial temperatures of $\\sim 10^4$ K.  We extract $\\sim 10^6$ ($10^7$) sightlines through the ionization field at $z=$ 6 (5) originating from identified QSO host halos.  However, we do not make detailed simulated spectra for this study.  We confirm that reionization is very inhomogeneous.  In fact, most sightlines through a $\\avenf\\lsim0.01$ universe do not intersect a single neutral region in a 500 Mpc stretch (corresponding to the redshift interval from $z=6$ to $z=5$. Furthermore, QSO host halos reside in very biased locations of the ionization field, where neutral patches are even rarer. Even without explicitly including the quasar's own ionizing radiation, sightlines originating from QSO halos are over twice as likely than those from random locations to entirely go through the ionized IGM within the first 60--80 Mpc, for $\\avenf\\lsim0.1$.  As measured by the number density of HI patches, the bias is less extreme, and the asymptotic values are approached at distances of $\\sim$ 30--60 Mpc away from the quasar.  Even before overlap, the absorption in the \\lya\\ forest is dominated by the residual neutral hydrogen inside the ionized IGM at these redshifts.  Detecting the additional dark spectral regions caused by HI patches pre-overlap is very challenging, and cannot be done from the existing \\lya\\ forest data at $z\\sim6$. At $z\\sim5$, where there are more available sightlines and the forest is less dark, constraining $\\avenf\\lsim0.1$ might be possible given a large dynamic range from very deep spectra and/or the \\lyb\\ forest.  To further quantify the imprint of an incomplete reionization, we also create dark gap distributions.  In order to be systematic, we separately examine the dark gap PDFs arising from the residual hydrogen in the ionized IGM, and those from just the HI patches pre-overlap.  We find that these two PDFs are likely to have different slopes, which can result in a ``knee'' feature in the combined PDF.  At $z=5$, large dark gaps preferentially come from the HI patches pre-overlap; however, at $z=6$, large dark gaps could instead preferentially come from the ionized IGM.  Unfortunately, averaging the PDFs over redshift bins might smear out these signals, since the gap distribution from the ionized IGM has a strong redshift dependence.  A larger dynamical range would aid in the discriminating power of the current sample of QSO spectra, in both the mean absorption statistic and the dark gap distributions.  An increased dynamic range might be available with new instruments such as the {\\it JWST} and the TMT, or through long-integration with existing instruments such as Keck HIRES \\citep{BRS07} and ESI, or with higher-order Lyman lines, such as \\lyb. In the case of the later, one must be careful in modeling the fluctuating absorption from the superimposed lower redshift \\lya\\ forest.    We propose using the fraction of pixels which are dark as a simple, nearly model-independent upper limit on $\\avenf$.  This simple statistic might be sufficient to constrain $\\avenf \\lsim 0.1$ at $z\\sim5$ using deep spectra of the current sample of $z>5$ quasars.  In conclusion, there is currently no direct evidence that reionization was complete by $z\\sim$ 5 -- 6.  We stress that we are not promoting such a late reionization scenario.  This paper is merely intended as a caution against over-interpreting high-redshift observations.  \\vskip+0.5in  We would like to thank Michael Strauss and David Spiegel for interesting conversations.  We are grateful to Xiaohui Fan and Michael Strauss for providing a list of $z>5$ QSOs.  We also thank Steven Furlanetto, Adam Lidz, Zoltan Haiman, and Renyue Cen for comments on a draft version of this paper.  Support for this work was provided by NASA through Hubble Fellowship grant HST-HF-51245.01-A, awarded by the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., for NASA, under contract NAS 5-26555."
    },
    "0910/0910.1578.txt": {
        "abstract": "{}{}{}{}{} % 5 {} token are mandatory  \\abstract % context heading (optional) % {} leave it empty if necessary {} % aims heading (mandatory) %{We present the results of the PhD Thesis carried out by L\\'opez-S\\'anchez (2006). %\t{We have performed a comprehensive multiwavelength analysis of a sample of 20 starburst galaxies that %show the presence of a substantial population of very young massive stars, most of them classified as Wolf-Rayet (WR) galaxies. %%The main aims are the study of the massive star formation and the stellar %%populations in these galaxies and the role that interactions with or between %%dwarf galaxies and/or low surface companion objects have in the triggering of the bursts. %In this series of papers we present new optical and near-infrared photometric and spectroscopic observations and complete with data at other %wavelengths (X-ray, far-infrared and radio) available in the literature. In this paper, the second of the series, %we present the results of the analysis of long-slit %intermediate-resolution spectroscopy of star-formation bursts for 16 galaxies of our sample.} %analyze the physical conditions, chemical abundances, and kinematics of the %ionized gas via long-slit intermediate-resolution spectroscopy.} {We have performed a comprehensive multiwavelength analysis of a sample of 20 starburst galaxies that show the presence of a substantial population of very young massive stars, most of them classified as Wolf-Rayet (WR) galaxies. In this paper, the second of the series, we present the results of the analysis of long-slit intermediate-resolution spectroscopy of star-formation bursts for 16 galaxies of our sample.} % methods heading (mandatory) %{We completed new long-slit intermediate-resolution spectroscopy of %the star-formation bursts within our sample of Wolf-Rayet galaxies. %  {We study the spatial localization of the WR stars in each galaxy via the detection of the broad and/or nebular \\HeII\\ $\\lambda$4686 %emission line. We analyze the excitation mechanism and derive the reddening coefficient, physical conditions and %%excitation, density, high and low ionization temperature, reddening and nature of the \\HII regions. %chemical abundances of the ionized gas. In the most of the cases, the chemical abundances were derived using a direct determination %of the electron temperature. %%The analysis of the chemical abundances in different emission knots or morphological structures of a given galaxy %%allows to explore the presence of co-spatial objects with different chemical evolution. %We study the kinematics of the ionized gas to check the %rotation/turbulence pattern of each system. When possible, tentative estimates of the Keplerian mass of the galaxies have been calculated.} {We study the spatial localization of the WR stars in each galaxy. We analyze the excitation mechanism and derive the reddening coefficient, %%@ physical conditions and chemical abundances of the ionized gas. We study the kinematics of the ionized gas to check the rotation/turbulence pattern of each system. When possible, tentative estimates of the Keplerian mass of the galaxies have been calculated.} % results heading (mandatory) %     {Aperture effects and the exact positioning of the slit onto the WR-rich bursts seem to play a fundamental role in their detection. %We checked that the ages of the last star-forming burst estimated using optical spectra agree with those derived from \\Ha\\ imagery. %%We confirmed that the main excitation mechanism of the ionized gas is photoionization by massive stars in all the objects. %The derived physical conditions and chemical abundances are usually in agreement with previous observations reported in the literature. %Our analysis has revealed that 14 up to 20 systems show clear evidences of perturbed kinematics; some kinematic anomalies are also found in other %4 galaxies. With respect to the results found in individual objects, we remark the detection of an knot with different metallicity an decoupled %kinematics in Haro~15, the finding of intense tidal streams in IRAS~08208+2816, UM~420, Tol~9 and perhaps in SBS~1319+579, and the confirmation of a %merging process in SBS~0926+606~A and in Tol~1457-262.} {Aperture effects and the exact positioning of the slit onto the WR-rich bursts seem to play a fundamental role in their detection. We checked that %%@ the ages of the last star-forming burst estimated using optical spectra agree with those derived from \\Ha\\ imagery. Our analysis has revealed that a substantial fraction of the galaxies show evidences of perturbed kinematics. With respect to the results found in %%@ individual galaxies, we remark the detection of objects with different metallicity and decoupled kinematics in Haro~15 and Mkn~1199, the finding of evidences of tidal streams in IRAS~08208+2816, Tol~9 and perhaps in SBS~1319+579, and the development of a merging process in SBS~0926+606~A and in Tol~1457-262.} % conclusions heading (optional), leave it empty if necessary {All these results --in combination with those obtained in Paper I-- reinforce the hypothesis that interactions with or between dwarf objects is a very important mechanism in the triggering of massive star formation in starburst galaxies, specially in dwarf ones. It must be highlighted that only deep and very detailed observations --as these presented in this paper-- can provide clear evidences that these subtle interaction processes are taking place.}   \\titlerunning{Massive star formation in Wolf-Rayet galaxies II: Spectroscopic results}  \\authorrunning{L\\'opez-S\\'anchez \\& Esteban} ",
        "introduction": "Wolf-Rayet (WR) stars are the evolved descendants of the most massive, very hot %(temperatures over $\\sim$50~000~K) and very luminous (10$^5$ to 10$^6$ \\Lo) O stars. In the so-called \\citet{Conti76} and \\citet{Maeder90,Maeder91} scenarios, WR stars are interpreted as central He-burning objects %%@ that have lost the main part of their H-rich envelope via strong winds. %(see van der Hucht 2001 for a review). Hence, their surface chemical composition is dominated by He rather than H, along with elements %like C, N and O, the produced by the nuclear nucleosynthesis. WN and WC stars show the products of the CNO cycle (H-burning) and the triple-$\\alpha$ (He-burning), respectively. The most massive O stars ($M\\geq$ \\mbox{25 \\Mo}\\ for \\Zo) became WR stars between 2 and 5 \\Myr\\ since their birth, spending only some few hundreds of thousands of years ($t_{WR}\\leq5\\times10^5$~yr) in this phase \\citep{MeynetMaeder05}. A review of the physical properties of WR stars was recently presented by %%@ \\citet{Crowther07}.  The broad emission features that characterized the spectra of WR stars are often observed in extragalactic \\HII regions. Actually, the so-called %%@ Wolf-Rayet galaxies make up a very inhomogeneous class of star-forming objects: giant \\ion{H}{ii} regions in spiral arms, irregular galaxies, blue %%@ compact dwarf galaxies (BCDGs), luminous merging \\IRAS\\ galaxies, active galactic nuclei (AGNs), Seyfert 2 and low-ionization nuclear emission-line %%@ regions (LINERs). All objects have in common ongoing or recent star formation which has produced stars massive enough to evolve to the WR stage %%@ (Shaerer et al. 1999). There are two important broad features that reveal the presence of WR stars in the integrated spectra of an extragalactic \\HII region: \\begin{enumerate} \\item A\tblend of \\ion{He}{ii} $\\lambda$4686, \\ion{C}{iii}/\\ion{C}{iv} $\\lambda$4650 and \\ion{N}{iii} $\\lambda$4640 emission lines originated in the expanding atmospheres of the most massive stars, the so-called {\\bf blue WR Bump}. This feature is mainly due to the presence of WN stars. The broad, stellar, \\ion{He}{ii} $\\lambda$4686 is its main feature. Although rarely strong, the narrow, nebular \\ion{He}{ii} $\\lambda$4686  is usually associated with the presence of these massive stars, but its origin is still somewhat controversial \\citep{Garnett91,G04,BKD08}. \\item The \\ion{C}{iii} $\\lambda$5698 and \\ion{C}{iv} $\\lambda$5808 broad emission lines, sometimes called the {\\bf red or yellow WR Bump}. %%@ \\ion{C}{iv} $\\lambda$5808 is the strongest emission line in WC stars but it is barely seen in WN stars. The red WR bump is rarely detected and it is always weaker than the blue WR bump \\citep{GIT00,FCCG04}. \\end{enumerate}  Making use of population synthesis models it is possible to determine the age of the bursts, the number of O and WR stars, the WN/WC ratio, the initial mass function (\\IMF) or the mass of the burst. Therefore, the study of WR galaxies helps to widen our knowledge about both the massive star formation and the evolution of starbursts: they allow to study the early phases of starbursts and are the best direct measure of the upper end of the \\IMF, a fundamental ingredient for studying unresolved stellar populations, \\citep{SGIT00,GIT00,PSGD02,FCCG04,Buckalew05,Zhang07}, giving key %%@ constraints to stellar evolution models.  %----------------  On the other hand, the knowledge of the chemical composition of galaxies, in particular in dwarf galaxies, is vital for understanding their %%@ evolution, star formation history, stellar nucleosynthesis, the importance of gas inflow and outflow and the enrichment of the intergalactic medium. %%@ Indeed, metallicity is a key ingredient for modelling galaxy properties, such it determines \\UV, optical and \\NIR\\ \\mbox{colors} at a given age %%@ (i.e., Leitherer et al. 1999), nucleosynthetic yields (e.g., Woosley \\& Weaver 1995), the dust-to-gas ratio (e.g., Hirashita et al 2001), the shape %%@ of the interstellar extinction curve (e.g., Piovan et al. 2006), or even the WR properties \\citep{Crowther07}.  The most robust method to derive the metallicity in star-forming and starburst galaxies is via the estimation of metal abundances and abundance %%@ ratios, in concrete through the determination of the gas-phase oxygen and nitrogen abundances and the nitrogen-to-oxygen ratio. The relationships %%@ between current metallicity and other galaxy parameters, such as colors, luminosity, neutral gas content, star formation rate, extinction or total %%@ mass, constraint galaxy evolution models and give clues about the current stage of a galaxy. For example, is still debated if massive star formation %%@ result in the instantaneous enrichment of the interstellar medium of a dwarf galaxy, or if the bulk of the newly synthesized heavy elements must cool %%@ before becoming part of the ISM that eventually will form the next generation of stars. Accurate oxygen abundance measurements of several \\HII %%@ regions within a dwarf galaxy will increase the understanding of its chemical enrichment and mixing of enriched material. The analysis of the %%@ kinematics of the ionized gas will also help to understand the dynamic stage of galaxies and reveal recent interaction features. Furthermore, %%@ detailed analysis of starburst galaxies in the nearby Universe are fundamental to interpret the observations of high-z star forming galaxies, such as %%@ Lyman Break Galaxies \\citep{EP03}, as well as quantify the importance of interactions in the triggering of the star-formation bursts, that seem to be %%@ very common at higher redshifts (i.e., Kauffmann \\& White 1993; Springer et al. 2005).  %-------------  We have performed a detailed photometric and spectroscopic analysis of a sample of 20 WR galaxies. Our main aim is the study of the formation of massive stars in starburst galaxies, their gas-phase metal abundance and its relationships with other galaxy properties, and the role that the %%@ interactions with or between dwarf galaxies and/or low surface brightness objects have in the triggering mechanism of the star-formation bursts. In Paper~I \\citep{LSE08} we exposed the motivation of this work, compiled the list of the %%@ analyzed WR galaxies (Table~1 of Paper~I) and presented the results of optical/\\NIR\\ broad-band and \\Ha\\ photometry. In this second paper we present %%@ the results of the analysis of intermediate-resolution long slit spectroscopy of 16 objects of our sample of WR galaxies --the results for the other 4 %%@ objects have been published separately. In many cases, two or more slit positions have been used in order to analyze the most interesting zones, knots or morphological structures belonging to each galaxy or even surrounding objects. In particular, these observations have the following aims: \\begin{enumerate} \\item Study the content and spatial location of the WR stars in each galaxy. We examine the spectra for the presence of the \\ion{He}{ii} %%@ $\\lambda$4686 emission line and/or the blue-\\WRBUMP\\ as well as for the red-\\WRBUMP. The characteristics of the WR population can be derived by %%@ comparison with theoretical population synthesis models. \\item Know the physical properties of the ionized gas: excitation mechanism, electron density, high and low ionization electron temperatures and %%@ reddening coefficient. \\item Analyze the ionization structure and the chemistry of the gas (abundances of He, O, N, S, Ne, Ar, Fe and Cl) associated with different morphological zones in each galaxy, especially in those areas in which WR features are detected. This analysis is specially relevant in cases of interaction or merging processes because the regions may have different chemical composition and also would allow to discern between the \\emph{tidal %%@ dwarf galaxy} (\\TDG) or \\emph{pre-existing dwarf galaxy} nature of nearby diffuse objects surrounding the main galaxy. \\item Determine the radial velocities of different star-formation bursts, galaxies in the same system and/or objects in possible interaction. The distance to the main galaxy is also calculated. \\item Study the velocity field via the analysis of position-velocity diagrams in order to understand the kinematics of the ionized gas associated to different members in the system in order to know their evolution (rotation, interactions features, fusion evidences, movements associated to superwinds...). The Keplerian mass has been estimated in objects showing solid-body rotation. \\item Obtain independent estimations of the age of the last star-forming burst via the comparison with stellar population synthesis models. \\item Study the stellar population underlying the bursts using the analysis of absorption lines (i.e. \\ion{Ca}{ii} H,K, \\ion{Mg}{i} $\\lambda\\lambda$5167,5184, \\ion{Na}{i} $\\lambda\\lambda$5890,5896, \\ion{Ca}{ii} triplet). \\item Finally, the spectral energy distribution (\\SED) has been analyzed in some cases in order to constrain the properties of the underlying stellar %%@ population. \\end{enumerate}   This paper mainly presents the analysis of the ionized gas within our WR galaxy sample. In \\S2 we describe our observations, some details of the data reduction processes, and describe some useful relations. In \\S3 we describe the %%@ physical properties ($T_e$, $n_e$, reddening coefficient, excitation mechanism), the chemical abundances and the kinematics of the ionized gas for %%@ each galaxy. Finally, the most important results derived from our spectroscopic study, including a comparison with the ages derived from the \\Ha\\ %%@ photometry and an estimation of the age of the underlying stellar component using the \\SED, are summarized in \\S4. A detailed analysis of the O and WR populations and the comparison with theoretical models is presented in Paper~III. The global analysis of our %%@ optical/NIR data will be shown in Paper IV. The final paper of the series (Paper~V) will compile the properties derived using data from other %%@ wavelengths (UV, FIR, radio and X-ray) and complete the global analysis combining all available multiwavelength data of our WR galaxy sample. It is, so far, the most complete and %%@ exhaustive data set of this kind of galaxies, involving multiwavelength results and analyzed following the same procedures. ",
        "conclusions": ""
    },
    "0910/0910.4944_arXiv.txt": {
        "abstract": "We present a catalog of 17 filamentary X-ray features located within a 68'$\\times$34' view centered on the Galactic Center region from images taken by \\textit{Chandra}.  These features are described by their morphological and spectral properties.  Many of the X-ray features have non-thermal spectra that are well fit by an absorbed power-law.  Of the 17 features, we find 6 that have not been previously detected, 4 of which are outside the immediate 20'$\\times$20' area centered on the GC. 7 of the 17 identified filaments have morphological and spectral properties expected for pulsar wind nebulae with X-ray luminosities of 5$\\times$10$^{32}$ to 10$^{34}$ ergs s$^{-1}$ in the 2.0-10.0 keV band and photon indexes in the range of $\\Gamma$=1.1 to 1.9. In one feature, we suggest the strong neutral Fe K$\\alpha$ emission line to be a possible indicator for past activity of Sgr A*. For G359.942-0.03, a particular filament of interest, we propose the model of a ram-pressure confined stellar wind bubble from a massive star to account for the morphology, spectral shape and 6.7 keV He-like Fe emission detected.  We also present a piecewise spectral analysis on two features of interest, G0.13-0.11 and G359.89-0.08, to further examine their physical interpretations.  This analysis favors the PWN scenario for these features. ",
        "introduction": "The Galactic Center (GC) region is a unique environment that is home to numerous energetic processes.  Radio, infrared and X-ray observations show structures that can only be distinguished in our own galaxy due to its proximity.  The GC thus provides a laboratory in which to probe the Galactic nuclear environment and the interactions between star formation regions, the interstellar medium (ISM) and the supermassive black hole Sgr A*.  These observations can aid in understanding the nuclear environments of nearby galaxies.  Radio observations have detected prominent non-thermal radio filaments that have been intensively studied \\citep[e.g.][]{Yusef84}. Many of the non-thermal radio filaments may come from milliGauss magnetic fields illuminated by energetic particles.  These fields are seen to be primarily perpendicular to the galactic plane and indicate the general structure of the magnetic field around the GC.  The exact origins and implications of the radio non-thermal filaments in relation to the GC are still in debate.  \\textit{Chandra} observations have detected numerous large scale X-ray structures from thermal to non-thermal filaments \\citep[e.g.][]{bamba02,lu08,muno08} primarily in the local 20'$\\times$20' area centered on the GC.  The commonly accepted scenarios for non-thermal thread-like X-ray features in the GC include pulsar wind nebulae (PWNe) \\citep[e.g.][]{wang93,wang06}, with the filament-like structure produced through ram-pressure confinement or strong magnetic fields, and magnetohydrodynamical shock fronts from supernova remnants (SNRs) \\citep[e.g][]{Yusef05}.  For the SNR case, the elongation of features represent shock fronts in the ISM.  The average non-thermal spectrum of a typical PWN is generally well fit by a power law model with a photon index in the range of $\\Gamma$=1.1-2.4 \\citep{gotthelf02} and X-ray luminosities in the range of $10^{32}$ to $10^{37}$ ergs s$^{-1}$ for the 0.2-10 keV energy band \\citep[e.g.][]{gaensler06,kaspi06}.  The PWN scenario is motivated primarily by the similarities between the features seen in the GC and X-ray features associated with known pulsars \\citep[see also][]{karga08}.  Few cases of K$\\alpha$ transitions from neutral iron (Fe) are proposed to stem from the reflection of hard X-rays from external sources, such as the supermassive black hole Sgr A* \\citep[e.g.][]{koyama89}.  By analyzing these X-ray features, one can in principle trace gas dynamics and magnetic fields in the GC region \\citep[e.g.][]{wang02} and examine the history of GC activity \\citep{koyama89}.  Previous broad scale studies of X-ray features focus on the immediate 20'$\\times$20' area surrounding Sgr A* \\citep[e.g.][]{muno08,lu08}. Here, we present a wide scale study on 17 X-ray thread-like features.  The features are detected over a roughly 68'$\\times$34' area of the sky, relating to approximately 158 pc by 79 pc, around Sgr A* based on all available observations.  While some of these features have been examined in previous literature, the improved counting statistics of this study allows for piecewise analysis that will be able to detect trends in the spectra and therefore better constrain the physical model \\citep[see also][]{lu08}.  In addition to the piecewise analysis of some features, we also present an additional physical model in order to explain a feature with a Helium (He)-like Fe K$\\alpha$ emission line, comet-like morphology and steep spectrum.  For consistency within this study, the spectral and morphological characteristics were performed for all features listed. All physical distance measurements assume an 8 kpc distance to the GC and all errors, unless otherwise stated, are given to the 90\\% confidence level. ",
        "conclusions": "To conclude, X-ray features in the GC likely have a heterogeneous origin. Non-thermal features are typically produced through synchrotron emission of PWNe or magnetohydrodynamical shock fronts from SNRs.  For the PWN model, one would expect elongation due to either the pulsar movement in the ISM or magnetic field confined flows of the pulsar wind materials.  In such a case, one would expect to see spectral softening across the direction of elongation along with typical photon indexes of $\\Gamma$=1.1-2.4.  For the SNR case, the direction of elongation would be caused by shock fronts and any measurable spectral change would be expected across the width of the feature.  Of the features presented, 7 exhibit characteristics of PWN, although the SNR case has not been completely ruled out.  Their non-thermal spectra as well as typical comet-like morphology are consistent with the ram-pressure confined PWN model.  The lack of any detected pulsar signal from these PWNe prevents a definitive conclusion on their origin.  Some of these pulsars may be detected in future pulsar searches using high frequency and high resolution radio observations.  In the case of G359.942-0.03, the model of a ram-pressure confined stellar wind bubble best explains the detection of the 6.7 keV He-like Fe emission and steep spectral shape. The near-IR images show an apparent counterpart of G359.942-0.03 consistent with a massive star where the lack of extended near-IR emission could be due to a combination of extinction and inefficiency in ionizing photon conversion.  In addition to their physical models, the X-ray features provide information into the underlying structure of the GC.  Features such as G0.017-0.044 and the \"clumps\" as studied by \\cite{lu08} map out the history of the GC (e.g. the possible radiative illumination of G0.017-0.044 from Sgr A*).  The orientations of PWN-like features help to map the magnetic, gas and stellar dynamics of the GC. \\cite{lu08} predicted that the orientation of comet like features is due to a strong galactic nuclear wind comparable to typical pulsar velocities oriented away from Sgr A*.  Of the features present within a 20'$\\times$20' image centered on Sgr A* (see Fig. 2), this effect can be seen in 5 out of the 7 PWN candidates (excluding G359.95-0.04 and G359.89-0.08) and G359.933-0.037.  Those features outside the 20'$\\times$20' region along with the features G359.964-0.03, G359.95-0.04 and G359.89-0.08 do not necessarily follow this claim; though this could be explained by variations in the vector field of the galactic wind as a result of regions of varying density. Additionally, the filament-like features could be regulated by local ordered magnetic field lines as with G0.13-0.11's morphology.  These features can thus be used in more in-depth studies of GC dynamics, beyond the quick analysis just presented, in combination with deep observations at other wavelengths; e.g. radio polarization and magnetic field tracers.  For the GC observations, few X-ray filaments were identified in regions outside the central 20'$\\times$20' region.  These regions have integrated exposure times significantly less than the central region. In addition, few of the identified features within the central 20'$\\times$20' region were found in initial observations with exposure times $<$200 ks.  Thus, those presented in this study may represent only a tip of the \"iceberg\" of X-ray features.  Indeed, some may be aligned such that they appear like point sources with their direction of elongation along the line of sight.  Based on the average star formation rate of the GC \\citep[$\\sim 10^{-7}$ M$_{\\sun}$ yr$^{-1}$ pc$^{-3}$,][]{figer04}, \\cite{lu08} predicted $\\sim$40 pulsars within 7 pc of the GC with ages $\\lesssim 3\\times 10^5$ yr assuming an average stellar mass of 10$M_{\\sun}$.  This is roughly consistent with our 13 PWN candidates if we consider that projection effects will limit us to $\\sim$71\\% of PWN, assuming the extended X-ray emission of a PWN is identified if the angle between the extended emission and the line of sight lies between $\\pi$/4 and 3$\\pi$/4, along with factors including sensitivity limits and pulsar population statistics.  With higher sensitivity and spatial resolution in multiple wavelength bands, including X-ray, radio and near-IR, the identities of these X-ray features can be more accurately determined."
    },
    "0910/0910.1085_arXiv.txt": {
        "abstract": "This study resolves a discrepancy in the abundance of Zr in the 47 Tucan\\ae\\ asymptotic giant branch star Lee~2525. This star was observed using the echelle spectrograph on the 2.3~m telescope at Siding Spring Observatory. The analysis was undertaken by calibrating Lee~2525 with respect to the standard giant star Arcturus. This work emphasises the importance of using a standard star with stellar parameters comparable to the star under analysis rather than a calibration with respect to the Sun \\citep{Koch2008}. Systematic errors in the analysis process are then minimised due to the similarity in atmospheric structure between the standard and programme stars. The abundances derived for Lee 2525 were found to be in general agreement with the \\cite{Brown1992} values except for Zr. In this study Zr has a similar enhancement ([Zr/Fe]~$=+0.51$~dex) to another light {\\it s}-process element, Y ([Y/Fe]~$=+0.53$~dex), which reflects current theory regarding the enrichment of {\\it s}-process elements by nuclear processes within AGB stars \\citep{Busso2001}. This is contrary to the results of \\cite{Brown1992} where Zr was under-abundant ([Zr/Fe]~$=-0.51$~dex) and Y was over-abundant ([Y/Fe]~$=+0.50$~dex) with respect to Fe. ",
        "introduction": "\\label{sec:Intro_BW} The globular cluster 47~Tucan\\ae\\ (47~Tuc) has proven to be a rich source of study with regards to the structure and evolution of stars and stellar systems. Chemical abundance studies have indicated the presence of light element abundance anomalies between the stars at all stages of stellar evolution. Recent advances in telescope and instrument sensitivity are providing observation of larger samples of stars within clusters at higher resolution enabling a more detailed investigation of the exact nature of the abundance anomalies. A significant advance has been the greater number of abundances derived for the elements heavier than iron. In particular the light and heavy {\\it s}-process elements which provide signatures of key stages in stellar evolution.  A well studied phenomenon in 47~Tuc is the CN weak, CN strong bimodality which is seen at all stages of stellar evolution in 47~Tuc \\citep{Cannon1998}. Internal mixing cannot solely explain this anomaly and some primordial or pollution mechanism is required to account for it being observed on the main sequence and giant branches (\\citealt{Cannon1998}; \\citealt{Briley2004}). Also observed in 47~Tuc is a correlation of Na to CN strength \\citep{Cottrell1981}. However there has been no observational evidence of a correlation of either Mg or Al with CN or Na as can be observed in more metal poor clusters though Na has been observed to anti-correlate with O in 47 Tuc \\citep{Carretta2004}. As these anomalies have been observed in giants and dwarfs it is theorised they are most likely due to primordial scenarios rather than solely due to internal mixing in globular cluster stars. For a complete summary of these anomalies in 47 Tuc and other globular clusters refer to the recent review paper, \\cite{Gratton2004}.  Recent stellar studies have investigated the abundances of elements heavier than iron and have shown an enhancement in {\\it s}-process elements in 47~Tuc. \\cite{Brown1992} analysed four giant stars in 47~Tuc for their light and heavy element abundances, determining for the {\\it s}-process elements an average enhancement in Y ([Y/Fe]~$=+0.48\\pm0.11$~dex) but a depletion in Zr ([Zr/Fe]~$=-0.22\\pm0.05$~dex). \\cite{Alves-Brito2005} analysed a sample of five 47~Tuc giant stars and found Zr to also be depleted ([Zr/Fe]~$=-0.17\\pm0.12$~dex) though did not obtain abundances for Y. However \\cite{Wylie2006} found an overall enhancement in {\\it s}-process elements, including Zr ([Zr/Fe]~$=+0.65\\pm0.16$~dex) and Y ([Y/Fe]~$=+0.64\\pm0.20$~dex), in seven giant branch stars in 47 Tuc. In a study of 3 turn-off and 8 subgiants in 47~Tuc \\cite{James2004b} found enhancements in Sr and Ba, but Y was depleted for the subgiants ([Y/Fe]~$=-0.11\\pm0.10$~dex) and slightly enhanced for the turnoff stars ([Y/Fe]~$=+0.06\\pm0.010$~dex). While the results are consistent within each study the variation between studies are contradictory to the expected homogeneous distribution of heavy elements in globular cluster stars.  The primary source for {\\it s}-process element production are nuclear reactions that occur during the asymptotic giant branch (AGB) stage of stellar evolution \\citep{Busso2001}. AGB stars of mass greater than 1.5~M$_{\\odot}$ undergo third dredge up (TDU) during which seed nuclei (Fe) are exposed to neutron fluxes. This builds up heavier and heavier nuclei resulting in the observed enhancements in {\\it s}-process element abundances in thermally pulsing AGB stars. Due to small reaction cross-sections, elements are first built up in the light {\\it s}-process peak ({\\it ls} = $\\langle$Sr, Y, Zr$\\rangle$) then in the heavy {\\it s}-process peak ({\\it hs} = $\\langle$Ba, La, Nd$\\rangle$) \\citep{Busso2001}. The ratios of [{\\it ls}/Fe], [{\\it hs}/Fe] and [{\\it hs/ls}] are used as indicators of the degree of {\\it s}-process enhancement and neutron flux. A depletion in Zr coinciding with an enrichment in Y, as determined in \\cite{Brown1992}, is not a likely result in a scenario of enrichment due to AGB stars. There are other potential sources for the {\\it s}-process elements. The AGB source is referred to as the `main' {\\it s}-process. The `weak' {\\it s}-process occurs during He-core burning in massive stars and can enhance light {\\it s}-process elements such as Sr and Y \\citep{Arlandini1999}. The {\\it r}-process, rapid neutron exposures typically in supernov\\ae, can also contribute to the abundances of {\\it s}-process elements. The contributions from these different sources can be disentangled by comparing heavy elements whose abundances are mainly {\\it s}-process (such as Y and Zr) to those who are mainly {\\it r}-process, such as Eu, to those who have contributions in different ratios from these sources \\citep{Travaglio2004}.  While there are these different potential sources for each of the {\\it s}-process elements, including contributions from the {\\it r}-process, for the most part it is argued that the abundance of a heavy element is homogeneous within a cluster. The differing results for Zr (and Y) need to be resolved in order to pursue the connections between heavy and light element abundances. These connections provide clues as to the nature of primordial and mixing scenarios that could lead to chemical abundance distributions that are seen in globular cluster stars. ",
        "conclusions": "Studies to date of {\\it s}-process element abundances in 47 Tuc stars have concluded that the observed {\\it s}-process element abundances are due to the primordial chemical composition of the cluster or some pollution event early in the cluster's history. While within each study there is agreement as to the magnitude of the {\\it s}-process element abundances, the abundances differ between the studies. The analysis of Lee 2525 carried out here resolves a discrepancy between \\cite{Brown1992} and \\cite{Wylie2006} as to the magnitude of the Zr abundance in this star. The current study found Lee 2525 to be enhanced in Zr at [Zr/Fe]~$=+0.51$~dex which is in agreement with the enhancement found for Y of [Y/Fe]~$=+0.53$~dex, another light {\\it s}-process element. This is in line with {\\it s}-process element enhancements found in 47 Tuc giant branch stars in \\cite{Wylie2006}. The Na-CN correlation reported in \\cite{Brown1992} for Lee 2525 was also found here, as well as an Na-O anti-correlation in line with other studies of light elements in 47 Tuc stars \\citep{Carretta2004}. These results support the premise that the abundance anomalies observed in 47 Tuc have a primordial or pollution-based origin.  While this study has resolved a discrepancy between two key papers the overall {\\it s}-process element abundance distribution within 47 Tuc is still not clear. It is necessary to expand the sample of 47 Tuc stars analysed for their {\\it s}-process element abundances in order to consolidate the results found here with those from other studies. This analysis was carried out differentially with respect to Arcturus in an effort to reduce systematic errors by calibrating the analysis process to a standard star of similar stellar parameters. This provides a solid framework for further study within this area."
    },
    "0910/0910.4496_arXiv.txt": {
        "abstract": "\\noindent We present new high resolution spectroscopy of the low mass X-ray binary Cyg X-2 which enables us to refine the orbital solution and rotational broadening of the donor star. In contrast with Elebert et al (2009) we find a good agreement with results reported in Casares et al. (1998). We measure $P=9.84450\\pm0.00019$ day, $K_2=86.5\\pm1.2$ km s$^{-1}$ and $V \\sin i=33.7\\pm0.9$ km s$^{-1}$. These values imply $q=M_{2}/M_{1}=0.34 \\pm 0.02$ and $M_{1}=1.71\\pm 0.21$ M$_{\\odot}$ (for $i=62.5 \\pm 4^{\\circ}$). Therefore, the neutron star in Cyg X-2 can be more massive than canonical. We also find no evidence for irradiation effects in our radial velocity curve which could explain the discrepancy between Elebert et al's and our $K_2$ values. ",
        "introduction": "Dynamical studies in low mass X-ray binaries (LMXBs hereafter) offer a promissing route to test the equation of state of nuclear matter. Soft equations of state are not stable above 1.6 M$_{\\odot}$ \\citep[e.g.][]{brown94} and hence finding a neutron star more massive than this limit would be a major advance in our knowledge of nuclear matter physics. LMXBs and, in particular, accreting millisecond pulsars (AMPs) are expected to harbour the most massive neutron stars because of the sustained accretion of matter during their long lifes \\citep{vandenheuvel95}. Unfortunately, dynamical studies are hampered by (1) the overwhelming accretion luminosity in persistent LMXBs (which swamps the donor star's spectrum) and (2) the extreme faintness of the companion star in transient AMPs during quiescence \\citep{davanzo09}.  A promissing route to overcome these limitations was proposed by \\cite{steeghs02}. X-ray irradiated donor stars can be betrayed by the detection of high excitation fluorescence NIII and CIII lines within the Bowen blend at $\\simeq$4634-50 \\AA . The Doppler shift of these narrow lines traces the orbit of the donor star and, therefore, can yield dynamical constraints in persistent LMXBs and transient AMPs during outburst. As a result of this strategy some good candidates for massive neutron stars have been found, namely X1822-371 \\citep{munoz05} and Aql X-1 \\citep{cornelisse07}.  In a few long period ($P>$ 1 day) systems, the donor is an evolved star which can be spectroscopically detected over the irradiated accretion disc. Cyg X-2 represents one of this exceptional cases. Its first orbital solution was presented by \\cite{cowley79} who show that the donor star is an A5-F2 III orbiting the neutron star in 9.843 days with a projected velocity $K_2=87\\pm3$ km s$^{-1}$. Almost 20 years later, this was refined by \\cite{casares98} (C98) who find an A9 donor with orbital parameters $P_{\\rm orb}=9.8444 \\pm 0.0003$ d and $K_2=88.0\\pm1.4$ km s$^{-1}$. This work also reports the first determination of the rotational broadening of the donor's absorption features ($V \\sin i = 34.2 \\pm 2.5$ km s$^{-1}$) which, in turn, implies a binary mass ratio $q=M_{2}/M_{1}=0.34 \\pm 0.04$. This value is remarkable as it implies a peculiar donor star, very undermassive for its spectral type and the probable outcome of an intermediate mass binary \\citep{king99, podsiadlowski00, kolb00}. But the implications for the compact object's mass are also remarkable. The orbital solution, combined with inclination constraints $i=62.5\\pm4^{\\circ}$ derived through ellipsoidal fits to UBV light curves gives a neutron star mass of 1.78 $\\pm$ 0.23 M$_{\\odot}$ \\citep{orosz99} and, hence, a good candidate for a massive neutron star. This result has been recently challenged by \\cite{elebert09} (E09) who measure $K_2=79\\pm3$ km s$^{-1}$ and hence $M_1$=1.5 $\\pm$ 0.3 M$_{\\odot}$.  Here we present new high resolution spectroscopic observations of Cyg X-2 with the main aim of refining the rotational broadening of the donor star and update the orbital parameters. In contrast with E09, we find a good agreement with previous results reported in C98 which give support to the presence of a massive neutron star in Cyg X-2. ",
        "conclusions": "We have revisited the determination of the system parameters in Cyg X-2 with 21 new high-resolution spectra obtained during 1999 and 2000. The new solution does not support the conclusions of E09 who claim for a significantly lower value for the radial velocity semiamplitude of the donor star. Instead, our results confirm previous determinations reported in C98. The discrepancy cannot be explained by X-ray irradiation because the sinusoidal shape of the radial velocity curve is not disturbed. In particular, we find (i) no evidence for orbital eccentricity and (ii) no significant deviations between two velocity points at phase $\\sim$ 0.65, when X-ray flux varies by a factor $\\sim$3. Our refined orbital parameters are $P=9.84450 \\pm 0.00019$ days, $K_2=86.5 \\pm 1.2$ km s$^{-1}$ and $V \\sin i = 33.7 \\pm 0.9$ km s$^{-1}$, which lead to $q=0.34\\pm0.02$, $M_{1} \\sin^{3} i  = 1.19 \\pm 0.06$ M$_{\\odot}$. Assuming $i=62.5 \\pm 4^{\\circ}$ from \\cite{orosz99} leads to $M_{1}=1.71\\pm 0.21$ M$_{\\odot}$ and, therefore, the possibility that Cyg X-2 harbours a neutron star more massive than canonical."
    },
    "0910/0910.3339_arXiv.txt": {
        "abstract": "We present far-infrared spectra, $\\lambda$=52--93 $\\mu$m, obtained with the {\\it Spitzer Space Telescope} in the Spectral Energy Distribution mode of its MIPS instrument, of a representative sample of the most luminous compact far-infrared sources in the Large Magellanic Cloud. These include carbon stars, OH/IR Asymptotic Giant Branch (AGB) stars, post-AGB objects and Planetary Nebulae, the R\\,CrB-type star HV\\,2671, the OH/IR red supergiants WOH\\,G064 and IRAS\\,05280$-$6910, the three B[e] stars IRAS\\,04530$-$6916, R\\,66 and R\\,126, the Wolf-Rayet star Brey\\,3a, the Luminous Blue Variable (LBV) R\\,71, the supernova remnant N\\,49, a large number of young stellar objects (YSOs), compact H\\,{\\sc ii} regions and molecular cores, and a background galaxy at a redshift $z\\simeq0.175$. We use the spectra to constrain the presence and temperature of cold dust and the excitation conditions and shocks within the neutral and ionized gas, in the circumstellar environments and interfaces with the surrounding interstellar medium (ISM). First, we introduce a spectral classification scheme. Then, we measure line strengths, dust temperatures, and IR luminosities. Objects associated with star formation are readily distinguished from evolved stars by their cold dust and/or fine-structure lines. Evolved stars, including the LBV R\\,71, lack cold dust except in some cases where we argue that this is swept-up ISM. This leads to an estimate of the duration of the prolific dust-producing phase (``superwind'') of several thousand years for both RSGs and massive AGB stars, with a similar fractional mass loss experienced despite the different masses. We tentatively detect line emission from neutral oxygen in the extreme RSG WOH\\,G064, which suggests a large dust-free cavity with implications for the wind driving. In N\\,49, the shock between the supernova ejecta and ISM is revealed in spectacular fashion by its strong [O\\,{\\sc i}] $\\lambda$63-$\\mu$m emission and possibly water vapour; we estimate that 0.2 M$_\\odot$ of ISM dust was swept up. On the other hand, some of the compact H\\,{\\sc ii} regions display pronounced [O\\,{\\sc iii}] $\\lambda$88-$\\mu$m emission. The efficiency of photo-electric heating in the interfaces of ionized gas and molecular clouds is estimated at 0.1--0.3\\%. We confirm earlier indications of a low nitrogen content in the LMC. Evidence for solid state emission features is found in both young and evolved objects, but the carriers of these features remain elusive; some of the YSOs are found to contain crystalline water ice. The spectra constitute a valuable resource for the planning and interpretation of observations with the {\\it Herschel Space Observatory} and the {\\it Stratospheric Observatory For Infrared Astronomy} (SOFIA). ",
        "introduction": "About the cycle of gas and dust that drives galaxy evolution, much can be learnt from the interfaces between the sources of feedback and the interstellar medium (ISM), and between the ISM and the dense cores of molecular clouds wherein new generations of stars may form. These regions are characterized by the cooling ejecta from evolved stars and supernovae, and clouds heated by the radiation and shocks from hot stars, in supernova remnants (SNRs) and young stellar objects (YSOs) embedded in molecular clouds.  The interaction regions with the ISM lend themselves particularly well to investigation in the infrared (IR) domain, notably in the 50--100 $\\mu$m region; cool dust ($\\sim20$--100 K) shines brightly at these wavelengths, and several strong atomic and ionic transitions of abundant elements (viz.\\ [O\\,{\\sc i}] at $\\lambda=63$ $\\mu$m, [O\\,{\\sc iii}] at $\\lambda=88$ $\\mu$m, and [N\\,{\\sc iii}] at $\\lambda=57$ $\\mu$m) provide both important diagnostics of the excitation conditions and a mechanism for cooling. These diagnostic signatures became widely accessible within the Milky Way, by virtue of the {\\it Kuiper Airborne Observatory} (KAO, see Erickson et al.\\ 1984) and the Long-Wavelength Spectrograph (LWS, Clegg et al.\\ 1996) onboard the {\\it Infrared Space Observatory} (ISO, Kessler et al.\\ 1996).  The gas-rich dwarf companions to the Milky Way, the Large and Small Magellanic Clouds (LMC and SMC), offer a unique opportunity for a global assessment of the feedback into the ISM and the conditions for star formation, something which is much more challenging to obtain for the Milky Way due to our position within it. The LMC and SMC are nearby ($d\\approx50$ and 60 kpc, respectively: Cioni et al.\\ 2000; Keller \\& Wood 2006) and already the scanning survey with the {\\it IR Astronomical Satellite} (IRAS, Neugebauer et al.\\ 1984) showed discrete sources of far-IR emission in them. The star-forming regions, massive and intermediate-mass stars, and ISM are also lower in metal content than similar components of the Galactic Disc, $Z_{\\rm LMC}\\approx0.4$ Z$_\\odot$ and $Z_{\\rm SMC}\\approx0.1$--0.2 Z$_\\odot$ (cf.\\ discussion in Maeder, Grebel \\& Mermilliod 1999). This offers the possibility to assess the effect metallicity has on the dust content and on the heating and cooling processes, and to study these in environments that are more similar to those prevailing in the early Universe than the available Galactic examples (cf.\\ Oliveira 2009).  The {\\it Spitzer Space Telescope} (Werner et al.\\ 2004) marries superb sensitivity with exquisite imaging quality, able to detect the far-IR emission from a significant fraction of the total populations of YSOs, massive red supergiants (RSGs), intermediate-mass Asymptotic Giant Branch (AGB) stars, Planetary Nebulae (PNe), and rare but extreme --- and important --- phases in the late evolution of massive stars such as Luminous Blue Variables (LBVs) and SNRs. The telescope also carried a facility, the MIPS-SED, to obtain spectra at 52--93 $\\mu$m, and we used this to target representative samples of luminous 70-$\\mu$m point sources in the LMC and SMC. The SMC spectra are presented in Paper II in this two-part series (van Loon et al.\\ 2009); here we present the results of the LMC observations. ",
        "conclusions": "We have presented the 52--93 $\\mu$m spectra of 48 compact far-IR sources in the Large Magellanic Cloud, obtained with MIPS-SED onboard the {\\it Spitzer Space Telescope} as part of the SAGE-Spec Legacy Program. The spectra were classified using a simple classification scheme introduced here for the first time. We measured the intensity of the fine-structure lines of oxygen --- [O\\,{\\sc i}] at 63 $\\mu$m and [O\\,{\\sc iii}] at 88 $\\mu$m --- as well as the slope of the dust continuum emission spectrum which we translated into a dust temperature and far-IR luminosity. Some of the most interesting results arising from this analysis may be summarized as follows:  \\begin{itemize} \\item[$\\bullet$]{Young Stellar Objects (YSOs) separate rather well from more evolved objects by the steeper slope of the far-IR spectrum (i.e.\\ lower dust temperature) and to some extent brighter fine-structure line emission. The contributions from [O\\,{\\sc i}] and [O\\,{\\sc iii}] show large variations among the sample, likely due to different excitation mechanisms at play.} \\item[$\\bullet$]{The supernova remnant N\\,49 displays spectacular [O\\,{\\sc i}] line emission, contributing 11\\% to the MIPS 70-$\\mu$m broad-band flux. We interpret this as arising from shocked gas at the interface of the expanding supernova ejecta and the local ISM. A possible detection of water vapour and/or OH emission is reported. From the dust continuum emission we estimate that 0.2 M$_\\odot$ of dust is associated, consistent with the expected amount of collected ISM.} \\item[$\\bullet$]{Equally spectacular emission is seen in the [O\\,{\\sc iii}] line in 30\\,Dor-17 and N\\,159-P2. These objects are associated with regions of active star formation and are probably YSOs harboring an (ultra-)compact H\\,{\\sc ii} region.} \\item[$\\bullet$]{The ratio of [O\\,{\\sc i}] to IR (dust-processed) luminosity is used to estimate the efficiency of photo-electric heating in the interfaces between ionized gas and molecular clouds in YSOs and compact H\\,{\\sc ii} regions; we find it is $\\approx0.1$--0.3\\%.} \\item[$\\bullet$]{We derive a low nitrogen content of H\\,{\\sc ii} regions in the LMC, with a nitrogen-to-oxygen ratio of $N({\\rm N})/N({\\rm O})<0.1$, possibly $<0.02$. This confirms the sparse evidence previously available from far-IR fine-structure lines.} \\item[$\\bullet$]{The [O\\,{\\sc i}] line is detected in the extreme red supergiant WOH\\,G064. It is strong compared to the line detected in Galactic red supergiants, hinting at a larger dust-free cavity at lower metallicity. The implication may be a larger contribution of Alfv\\'en waves to the wind driving at lower metallicity.} \\item[$\\bullet$]{The circumstellar envelopes of the OH/IR stars in our sample lack cold dust. We estimate that dust production must have increased dramatically only a few thousands years ago, for the red supergiants and massive AGB stars alike. We further estimate that the {\\it fractional} mass lost in this ``superwind'' phase is rather similar for these stars in spite of their different birth masses.} \\item[$\\bullet$]{The Luminous Blue Variable R\\,71 also lacks cold dust, but the timescales are somewhat longer and the coolest dust may have formed a few $10^4$ yr ago; this would be consistent with an origin in the wind of a red supergiant progenitor to R\\,71, as suggested previously by others.} \\item[$\\bullet$]{In the case of the carbon stars in our sample, which show far-IR excess emission over that expected from a circumstellar shell, and possibly also some planetary nebulae, we interpret the far-IR emission as arising from swept-up ISM and not due to dust produced by the stars themselves.} \\item[$\\bullet$]{Broad (several $\\mu$m wide) emission features are seen in many sources, both young and evolved. It is likely that these arise from modes within solid state material, probably minerals, but the exact carriers could not be identified (maybe hydrous silicates).} \\item[$\\bullet$]{Emission in the 60--70 $\\mu$m region due to crystalline water ice is detected tentatively in several YSOs.} \\end{itemize}"
    },
    "0910/0910.5035_arXiv.txt": {
        "abstract": "B2352+495 is a prototypical example of a Compact Symmetric Object (CSO). It has a double radio lobe symmetrically located with respect to a central flat spectrum radio core (the location of the AGN) and has a physical extent of less than 200$\\,$pc. In this work we report VLBA observation of 21$\\,$cm \\HI~absorption toward B2352+495 to investigate the properties of this remarkable radio source, in particular, to explore whether the radio emission can be confined by circumnuclear material (frustration scenario) or whether the source is likely to be young. We confirmed the two \\HI~absorption features previously detected toward B2352+495 -- a broad line nearly centered at the systemic velocity of the galaxy and a narrow redshifted component. The atomic gas from the broad absorption component is likely associated with circumnuclear material, consistent with the current paradigm of clumpy \\HI~distribution in toroidal structures around supermassive black holes. ",
        "introduction": "Compact Symmetric Objects (CSOs) are characterized by compact bright radio emission at scales $<1\\,$kpc with lobes on both sides of a central engine (Wilkinson et al.$\\,$1994). CSOs tend to be associated with bright elliptical galaxies, although some appear to be associated with quasars and even spiral galaxies (Readhead et al.$\\,$1996; Augusto et al.$\\,$2006; Perlman et al.$\\,$1996). Many CSOs have an overall GigaHertz-Peaked Spectrum (GPS), and the cores are often undetected at frequencies of 5$\\,$GHz or higher (Taylor et al.$\\,$1996). CSOs are very luminous radio sources, significantly brighter (by more than a factor of 10) than the luminosity threshold between FR$\\,$I and II radio galaxies (i.e., $\\sim 2\\times 10^{25}\\,$W$\\,$Hz$^{-1}$sr$^{-1}$ at 178$\\,$MHz, Fanaroff \\& Riley 1974; Readhead et al.$\\,$1996).  It has been proposed that the small source size could be caused by ``frustration'' of the radio jets, which fail to escape from a dense galactic medium (e.g., van Breugel et al.$\\,$1984; O'Dea et al.$\\,$1991; De Young 1993). However, a different scenario is currently preferred: CSOs may be young progenitors ($\\la$10$^3$ -- 10$^5\\,$years) of powerful extended radio sources (e.g., Phillips \\& Mutel 1982; Augusto et al.$\\,$2006; Augusto 2007). Based on XMM-Newton observations of five CSOs, Vink et al.$\\,$(2006) found that their absorption column density ($N_H \\sim 10^{22}\\,$cm$^{-2}$) is on average greater than that of broad-line radio galaxies, but similar to column densities observed toward narrow-line and other types of radio galaxies (NLRGs, RG; see Sambruna et al.$\\,$1999 for definitions). These results argue against the frustration scenario because much higher X-ray absorption column densities than found by Vink et al. (2006) are expected if the expansion of the radio jets were confined (see also Guainazzi et al. 2006). In contrast, the dynamical age of CSOs measured from the expansion velocity of hotspots in the radio lobes supports the youth hypothesis (Taylor et al.$\\,$2000, Polatidis \\& Conway 2003). The youth interpretation is also supported by limits on molecular gas content (O'Dea et al.$\\,$2005), \\HI~observations (Pihlstr\\\"om et al.$\\,$2003), and under-luminous O{$\\,$\\small III} line emission (Vink et al.$\\,$2006).  The study of CSOs is therefore fundamental for the understanding of the formation and evolution of radio galaxies. In addition, CSOs are good candidates to investigate the fueling of supermassive black holes given that the compact radio emission enables very high angular resolution observations of circumnuclear gas via absorption studies. In this work we report high sensitivity and high angular resolution observations of \\HI~toward B2352+495; an archetype of the CSO class.   \\subsection{\\HI~Absorption in Compact Symmetric Objects}  According to Curran et al.$\\,$(2008), \\HI~absorption associated with sources at redshifts greater than $\\sim 0.1$ has been detected toward 37 sources, 13 of them classified as CSOs, GPS, or HFP (High Frequency Peaked) sources. An overview of \\HI~absorption in compact radio sources was published by Vermeulen et al.$\\,$(2003a). They report an \\HI~detection rate of $33\\%$ based on Westerbork Synthesis Radio Telescope (WSRT) observations of 59 sources (i.e., 19 detections). Of the 59 sources, 11 were classified as CSO by Augusto et al.$\\,$(2006) and Polatidis \\& Conway (2003). From this subsample of 11 CSOs, 6 (54$\\%$) show \\HI~absorption. Other studies show similar results, e.g., Peck \\& Taylor (2001) report a CSO \\HI~detection rate of $\\sim 50\\%$, and Gupta et al.$\\,$(2006) report a detection rate of $\\sim 45\\%$ in GPS sources (see also Vermeulen 2002).  The optical depth range of the 6 CSOs with \\HI~absorption from Vermeulen et al.$\\,$(2003a) is 0.2 to 1.7$\\%$, with linewidths between 13 and 297\\kms~and column densities between $2.8\\times 10^{19}$ and $2.6\\times 10^{20} \\,$cm$^{-2}$. To estimate \\HI~column densities, Vermeulen et al.$\\,$(2003a) assumed that the atomic gas was uniformly covering the radio continuum (covering factor of 1) and a spin temperature of 100$\\,$K, that may underestimate the spin temperature by more than an order of magnitude if the gas were located at parsec-scales from the AGN. Thus, the column densities reported by Vermeulen et al.$\\,$(2003a) are lower limits.  Including the results reported in this work, \\HI~imaging at high angular resolution has been reported toward five CSOs (Table~1). We also list in Table~1 all non-CSO AGN and starbursts that have been imaged at high angular resolution in \\HI, i.e., in interferometric observations using baselines greater than 100$\\,$km. In the rest of this section we review the most salient characteristics of the CSOs that have been studied in \\HI~at high angular resolution:   \\noindent {\\it --- B1946+708:} High angular resolution \\HI~observations were conducted by Peck et al.$\\,$(1999; see also Peck \\& Taylor 2001). They found extended and multi-peaked absorption throughout the radio continuum source. The peak optical depths range between 3 and 7$\\%$. The maximum column density ($N_{HI} \\sim 3 \\times 10^{23}\\,$cm$^{-2}$; Peck \\& Taylor 2001) was found toward the AGN core, which also shows the greatest velocity dispersion ($FWHM \\sim 350$\\kms). The extent of the \\HI~absorption in B1946+708 implies that the atomic gas is tracing a thick torus (thickness $> 80\\,$pc).  \\noindent {\\it --- 4C$\\,$31.04:} The host galaxy is a bright elliptical at z = 0.0592. It was classified as CSO based on VLBI continuum observations that revealed a double radio lobe with a flat spectrum radio core (Cotton et al.$\\,$1995; Giovannini et al.$\\,$ 2001, Giroletti et al.$\\,$2003; see also Augusto et al.$\\,$1998). Detection of \\HI~in 4C$\\,$31.04 was reported by Mirabel (1990) based on Arecibo observations. Mirabel (1990) found two absorption features, a broad (FWHM = 133\\kms) component and a double peaked narrow component (FWHM = 6 and 16\\kms) redshifted approximately 400\\kms~from the optical velocity. Subsequent VLBI observations by Conway (1999) showed that the broad line covers one of the radio lobes and just partially the second, indicating a sharp edge of the atomic gas torus. The narrow component is detected toward both radio lobes. The maximum optical depth of the lines is $\\sim 5$ to $7\\%$. Mirabel (1990) interpreted the narrow absorption as being caused by clouds similar to high velocity \\HI~clouds in our Galaxy.  \\noindent {\\it --- 4C$\\,$37.11:} In contrast to the two previous CSOs, \\HI~in 4C$\\,$37.11 has been detected only toward one of the radio lobes. The absorption is characterized by two broad (FWHM$>$100\\kms) \\HI~features from two separated (in line-of-sight and velocity) clouds. The peak optical depth of the \\HI~lines is $\\sim 2\\%$. Rodr\\'{\\i}guez et al.$\\,$(2009) concluded that the \\HI~absorption likely originates in a thick circumnuclear torus.  \\noindent {\\it --- PKS$\\,1413+135$:} As in 4C$\\,$37.11, \\HI~absorption is detected only toward one radio lobe (Perlman et al.$\\,$2002). The \\HI~optical depth is $\\tau \\approx 1$ and the line is quite narrow, between 16 and 18\\kms. Perlman et al.$\\,$(2002) concluded that the \\HI~absorption is likely from a giant molecular cloud in the outer disk of the host galaxy. Although PKS$\\,1413+135$ was classified by Perlman et al.$\\,$(1996) and Gugliucci et al.$\\,$(2005) as a CSO, it has several characteristics that differ form a typical CSO. In particular, it was originally classified as a BL Lac object based on its optical spectrum (Bregman et al. 1981), the nucleus dominates by more than 50$\\%$ of the total cm flux density, the radio emission is variable, and the host appears to be a spiral galaxy (see Perlman et al.$\\,$1996).  Finally, two additional sources (B1934$-$638 and 4C$\\,$12.50) are worth discussing in detail. B1934$-$638 is a compact-double GPS source that has been classified as a CSO (Augusto et al. 2006), however, no flat-spectrum radio core has been detected (Tzioumis et al. 2002). Even though no \\HI~imaging has been reported at high angular resolution, the \\HI~absorption is against the compact radio continuum as shown by detection of absorption with three baselines of the LBA (V\\'eron-Cetty et al. 2000). The absorption is narrow (FWHM = 18\\kms) and redshifted 260\\kms~from the optical systemic velocity, thus, the \\HI~is probably not located in the nuclear region of the galaxy. We include B1934$-$638 in Table~1 although no VLBI \\HI~imaging is available and the radio core has not been detected.  In the case of 4C$\\,$12.50, high angular resolution \\HI~absorption observations were reported by Morganti et al.$\\,$(2004). They found \\HI~absorption against a weak counterjet and concluded that the \\HI~likely traces a jet/cloud interaction and not a circumnuclear torus or disk. This source was classified as a CSO by Lister et al.$\\,$(2003) despite several atypical characteristics. For instance, while the 1.3$\\,$GHz continuum image reveals a S-shaped distribution with a total extend of $\\sim 300\\,$pc (Morganti et al.$\\,$2004), the 15$\\,$GHz VLBA observations are consistent with a core-jet morphology plus a weak counterjet (i.e., inconsistent with the typical symmetrical morphology of CSOs). 4C$\\,$12.50 has no clear hotspots, it is classified as a CSS (Compact Steep Spectrum), the radio jet exhibits superluminal motions, some isolated features have unusually high ($\\sim 60\\%$) linear polarization in comparison to other CSOs, and more importantly, there appears to be continuity between the $\\sim 300\\,$pc compact radio jets and large scale ($>10\\,$kpc) radio lobes (Stanghellini et al.$\\,$2005). Therefore, we do not consider 4C$\\,$12.50 a bonafide member of the CSO class (see Table~1).   \\subsection{B2352+495: An Archetype Compact Symmetric Object}   B2352+495 has been studied in detail by a number of authors (e.g., Readhead et al.$\\,$1996, Taylor et al.$\\,$1996); in this section we summarize the characteristics of this object:  \\noindent {\\it Radio Continuum:} The most salient property of B2352+495 is its highly symmetric morphology at radio wavelengths. It exhibits an S-shaped morphology with a total extent of $\\sim 120\\,$pc (Readhead et al.$\\,$1996). High angular resolution radio continuum observations from 610$\\,$MHz to 15$\\,$GHz reveal symmetric hotspots and lobes; remarkably similar to the morphology of FR II objects (but at much smaller scale). The dynamical age of the system is $\\sim 1200\\,$yr based on proper motion studies of the hotspots in the radio lobes (Taylor et al.$\\,$2000). Taylor et al.$\\,$(1996) detected the radio core of the system, i.e., the putative location of the central engine. The core is approximately located symmetrically between the radio lobes, and drives a collimated jet that is brighter toward the northern side of the core. The radio jet is detected up to the brightest radio emission in the system (B1 and B2; Taylor et al 1996; see also $\\S 3$) which is also located between the two radio lobes. The hotspots in the radio lobes show no superluminal motion (Conway et al.$\\,$1992). Multi-frequency VLA observations give fractional polarization limits of $<1\\%$ for this source (Rudnick \\& Jones 1983). The radio spectral energy distribution of B2352+495 peaks at $\\sim 1\\,$GHz, thus by definition, it is a GPS object (Conway et al.$\\,$1992). The total flux density of B2352+495 shows only modest ($\\la 20\\%$) variability on time scales of months; no systematic long term variability was found over a period of $\\sim 15\\,$years (Waltman et al.$\\,$1991; Conway et al.$\\,$1992). The radio luminosity is $\\sim 10^{26}\\,h^{-2}\\,$W$\\,$Hz$^{-1}$ ($\\nu_o = $5$\\,$GHz at the emitted frame), i.e., some 20 times brighter than the luminosity boundary between FR Is and FR IIs (Readhead et al.$\\,$1996).   \\noindent {\\it Optical/Infrared Counterpart:} B2352+495 was not detected by IRAS to sensitivity limits well below the corresponding luminosity of Arp$\\,$220 (at the redshifted frequency), thus B2352+495 is significantly less luminous than ultra-luminous infrared galaxies (Readhead et al.$\\,$1996). Snellen et al.$\\,$(2003) conducted imaging and spectroscopic optical observations of B2352+495. They found that the radio source resides in an optical host in the fundamental plane of elliptical galaxies, i.e., the host is a normal elliptical.$\\,$Snellen et al.$\\,$(2003) report a redshift of 0.23790$\\pm$0.00016, stellar velocity dispersion of 201$\\pm$17\\kms, and an effective radius of 1.6\\arcsec$\\pm$0.3\\arcsec~(5.8$\\pm$1.1$\\,$kpc).\\footnote{At a redshift of 0.23790, 10$\\,$mas correspond to 36.3$\\,$pc assuming a cosmology of $H_0 = 73.0$\\kms$\\,$Mpc$^{-1}$, $\\Omega_{matter} = 0.27$, $\\Omega_{vacuum} = 0.73$.} Based on the stellar velocity dispersion, the mass of the supermassive black hole at the center of B2352+495 is of the order of $10^8$\\Mo. The absolute magnitude of the galaxy is $M_V = -19.8$ with an apparent magnitude of $m_V = 20.1$ (Readhead et al.$\\,$1996). The optical luminosity of B2352+495 is $\\sim 0.35\\,L^*$, which is smaller than the typical optical luminosity of Classical Doubles (FR IIs; $L \\sim L^*$; Owen \\& White 1991) and cD galaxies ($L \\sim 5 - 10\\,L^*$; Schechter 1976). The absence of Balmer lines and strong blue continuum in the spectrum demonstrates no starburst activity. In contrast, a number of emission lines are detected typical of active elliptical galaxies (Readhead et al.$\\,$1996). The AGN contribution to the total optical continuum is $\\la 10\\%$ (Snellen et al.$\\,$2003).   \\noindent {\\it X-Ray Properties:} High sensitivity XMM-Newton observations of B2352+495 were conducted by Vink et al.$\\,$(2006). They detected an X-ray source with a luminosity of $4.6 \\times 10^{42}\\,$erg$\\,$s$^{-1}$ ($2 - 10\\,$keV band, ignoring absorption). Based on the X-ray detection, they estimated an intrinsic absorption column density of (0.66$\\pm$0.27)$\\times 10^{22}\\,$cm$^{-2}$, which is similar to the column densities found toward extended radio galaxies.   \\noindent {\\it Prior \\HI~Observations: } Absorption of the \\HI~transition in B2352+495 was first detected by Vermeulen et al.$\\,$(2003a) based on WSRT observations. The \\HI~absorption in B2352+495 consists of a broad line (FWHM = 82\\kms) and a redshifted narrow component (FWHM = 13\\kms). The peak flux density of the two absorption components is similar ($\\sim 40\\,$mJy). Preliminary results of the data reported in this paper were discussed by Vermeulen (2002). ",
        "conclusions": "As in the case of the \\HI~gas associated with the nucleus of B1946+708 (Peck et al.$\\,$1999), we estimate the \\HI~column density assuming a covering factor of 1 and a spin temperature of 8000$\\,$K, i.e., assuming that the \\HI~gas is directly associated with the AGN. Using the standard column density/opacity relation (e.g., Rohlfs \\& Wilson 2000), we estimate a total column density of $(5.2\\pm0.2) \\times 10^{22}$ and $(7.3 \\pm 1.0) \\times 10^{21}\\,$cm$^{-2}$ for the broad and narrow \\HI~lines, respectively (only statistical errors from the fit were used to estimate the column density uncertainty). The values are greater than the B2352+495 \\HI~column densities reported by Vermeulen et al.$\\,$(2003a), however, as mentioned above, their values are lower limits because of the assumption of uniform coverage of the radio continuum (which is not the case, see Figure~1) and the assumption of a lower spin temperature ($T_{sp} = 100\\,$K). If the \\HI~clouds have physical dimensions comparable to the synthesized beam (note that there appears to be structure in the opacity channels, Figure~2), then the \\HI~density would be between $\\sim 10^2$ and 10$^3\\,$cm$^{-3}$, corresponding to total \\HI~masses of $\\sim 2\\times10^5$ and $3 \\times 10^4\\,$\\Mo~for the broad and narrow features, respectively.  If the \\HI~absorption profile of the broad component is the superposition of many narrow \\HI~features, and assuming a typical optical depth of $\\tau = 0.04$ for the narrow features, then we could not have detected \\HI~absorption (at a $\\ge$3$\\sigma$ level) anywhere except toward the brightest continuum region at the center of the CSO (i.e., the region where \\HI~absorption was indeed detected). This implies that, given our sensitivity limit, our data do not set significant constraints on the scale height of the \\HI~distribution.   As discussed below, we favor the interpretation that the broad \\HI~component is associated with rotation of a disk/torus around the supermassive black hole in B2352+495. The detection of a narrow \\HI~feature redshifted with respect to the systemic velocity indicates infall. However, we cannot establish whether the \\HI~feature is associated with circumnuclear material or with a foreground \\HI~cloud in the host galaxy, e.g, whether it is located in the narrow-line region (NLR; e.g., B2050+364, Vermeulen et al.$\\,$2006) or associated with a high velocity cloud. We note that other studies (e.g., van Gorkom et al.$\\,$1989, Peck \\& Taylor 1998, Conway 1999) have also found evidence for infalling material in galaxies with compact radio cores (see however Vermeulen et al.$\\,$2003a).  We now discuss the nature of the radio continuum having the \\HI~absorption. In principle, the brightest 1.14$\\,$GHz continuum source (located close to the center of the two radio lobes, Figure~1) is suggestive of the active core, i.e., the location of the supermassive black hole. However, comparing our continuum data with higher angular resolution observations reveal a different scenario. In Figure~3 (upper panel) we show our 1.14$\\,$GHz continuum map superimposed on the 15.4$\\,$GHz observations of Taylor et al.$\\,$(1996). Both continuum data sets were self-calibrated, and thus no absolute astrometry is available. To overlay the two images, we used an average offset computed from the peak in the northern and the southern radio lobe emission; we assume that position shifts due to opacity effects are small there, since their measured separation is the same within the errors at 1.14$\\,$GHz and 15.4$\\,$GHz. The brightest radio emission at the center of the radio lobes was {\\it not} used to estimate the offset.  Taylor et al.$\\,$(1996) found a compact radio source characterized by an inverted spectral index ($\\alpha = 0.14\\pm0.3$; $S_\\nu \\propto \\nu^\\alpha$), which convincingly indicates the location of the active core (see Figure~3). Assuming a flat spectral index between 1 and 8$\\,$GHz, the flux density of the core at 1$\\,$GHz is $\\sim 13\\,$mJy. Our \\HI~optical depth limit toward the radio core is $<$0.1 (3$\\sigma$). A prominent radio jet is detected to the north of the core, and a weaker counterjet to the south (Taylor et al.$\\,$1996; 2000). The northern edge of the jet is coincident with the brightest 1.14$\\,$GHz continuum source reported in this work, implying that the \\HI~absorption is not against continuum from the core, but radio emission from the jet. Among other possibilities, the detection of the radio jet and weaker counterjet could be due to relativistic beaming (see Taylor et al.$\\,$[1996] for this and other interpretations). The \\HI~absorption in B2352+495 is thus similar to the absorption observed in B0402+379 (Rodr\\'{\\i}guez et al.$\\,$2009) and NGC$\\,$4151 (Mundell et al.$\\,$1995, 2003) where \\HI~is detected only toward one side of the radio jet.  It is interesting to note that the radio continuum source toward which the \\HI~is detected (Figure~3) is the location where the radio jet begins to change its position angle. This would suggest that, as for example in the case of 3C$\\,$236 (Conway 1999), the \\HI~absorption marks the location of a jet/cloud interaction responsible for disturbing (and possibly bending) the jet. However, the B2352$+$495 morphology is quite symmetric in both radio lobes which suggests that the deviation of the radio jets is not due to local jet/cloud interactions but governed by the large scale magnetohydrodynamical properties of the galactic nucleus. In the next section, we propose a different interpretation for the nature of the \\HI~absorption in B2352+495.   \\subsection{Possible \\HI~Velocity Gradient and Disk/Torus Interpretation}  It is difficult to reliably investigate the velocity field of the broad \\HI~line given that the absorption is only detected toward the central continuum source and, as shown in Figure~2, the velocity field is complex. Nevertheless, a velocity gradient may be present: the two channels with stronger absorption at velocities above 10\\kms~from systemic appear to have stronger absorption toward the west, whereas most of the blueshifted absorption is stronger toward the east (velocity channels $< -10$\\kms~from systemic, see Figure~2). Figure~4 shows the position-velocity diagram of the broad \\HI~absorption line in the east-west direction. As also seen in the individual channel maps, there is a hint of a velocity gradient between blueshifted and redshifted material. To illustrate the possible gradient, we show in Figure~3 (top panel) the opacity channels that correspond to the rest frame velocities of 17.6 and $-31.4$\\kms~(Fig.~2).  The existence of a velocity gradient is tantalizing given that it is approximately centered at the systemic velocity of the galaxy, and the position angle of the gradient is almost perpendicular to the north-south jet, i.e., the \\HI~gas could be tracing some kind of rotating structure around the supermassive black hole. We explore this hypothesis in the rest of this section, however, we stress that more sensitive global-VLBI observations are required to confirm the velocity gradient.  To investigate whether rotation around the supermassive back hole is physically plausible, we followed the formulation presented by Rodr\\'{\\i}guez et al.$\\,$(2009; see their appendix). To model the \\HI~rotating structure, we used the position and velocity of the peak optical depth of the 17.6 and $-$31.4\\kms~channels (Figure~2). The position of the radio core (Figure~3 upper panel; see also Taylor et al.$\\,$1996) was assumed to be coincident with the center of mass of the system. For a black hole mass of 10$^8$\\Mo~(Snellen et al.$\\,$2003) and assuming circular trajectories coplanar with the supermassive black hole, we obtain that the orbit of the \\HI~clouds from the 17.6 and $-$31.4\\kms~channels would have a radius of $\\sim 20\\,$pc with an orbital inclination of $\\sim 65$\\arcdeg~with respect to the plane of the sky. Such an orbit is reasonable given the expected distribution of atomic clouds around AGNs (e.g., Peck 2005). Similar results are obtained in the case that the orbital plane of the \\HI~clouds is not coplanar with the supermassive black hole but located above by distances of the order of the projected separation between the AGN core and the \\HI~absorption. This estimate is clearly a zero-order approximation because non-radial inhomogeneities in the gravitational field due to stars and gas are neglected. In addition, the radius of influence of the supermassive black hole ($r_{BH} = G M_{BH} / \\sigma_*^2$, where $G$ is the gravitational constant, $M_{BH}$ is the mass of the black hole, and $\\sigma_*$ is the stellar velocity dispersion; e.g., Neumayer et al. 2007) is $\\sim 10\\,$pc, thus, the stellar mass enclosed by the atomic gas is likely comparable to the black hole mass.   In Figure~3 (bottom panel) we show two possibilities for the nature of the \\HI~distribution, i.e., a thick torus (characterized by an \\HI~scale height similar or greater than the length of the radio source) or a thin (most likely warped, e.g., Pringle 1996) disk. The obvious method to discriminate between these possibilities is to check whether there is \\HI~toward the southern radio lobe. However, as mentioned above, our optical depth sensitivity limit does not set significant constraints on the \\HI~distribution, i.e., clouds like those detected toward the brightest radio continuum emission could be present throughout the region.  An indirect approach to discriminate between the models is to compare our observations with lower angular resolution data. As mentioned in $\\S 2$, the peak flux density measured with WSRT (Vermeulen et al.$\\,$2003a) is approximately the same peak intensity we measured with the VLBA (Figure~1) and the linewidths are also consistent. Thus, in contrast to other systems where VLBI observations resolve/filter out significant \\HI~signal (e.g., IC$\\,$5063, Oosterloo et al.$\\,$2000), the \\HI~in B2352+495 is compact and mostly distributed along the central continuum source, i.e., consistent with the thin disk interpretation.   The thin disk interpretation is also circumstantially supported by X-ray observations. Vink et al.$\\,$(2006) reported a X-ray absorption column density of $(6.6\\pm2.7) \\times 10^{21}\\,$cm$^{-2}$, which is similar to the column density of the narrow \\HI~component, but several times smaller than the value of the broad line. If the \\HI~absorption originates from an inclined (and warped) thin disk that does not substantially obscure the AGN (Figure~3, bottom panel; see also Fig. 5 of Orienti et al.$\\,$2006) then as observed, the column density toward the AGN derived from X-ray observations is expected to be smaller than that of the neutral disk.  We note that the thin disk interpretation does not necessarily imply a misalignment between the jet and the disk, i.e., the position angle of the jet changes as a function of distance from the core (Figure~3) and the disk could be warped, thus, the jet and disk could be perpendicular at (sub)parsec scales. In addition, given that the detection rate of \\HI~absorption in CSOs is $\\sim 50\\%$ (see $\\S 1.1$), it is certainly possible that in some cases the atomic gas is distributed in a circumnuclear structure that is seen in projection only toward one radio lobe/jet.  Assuming the orbital parameters from the thin disk model, we now explore the stability of the \\HI~disk with respect to gravitational fragmentation. We used the Toomre stability formulation as presented by Gallimore et al.$\\,$(1999), i.e.,  \\begin{equation} \\frac{\\Sigma_{HI}}{\\Sigma_c} \\simeq \\frac{0.028 N_{HI} v_{rot}}{r_{pc}}, \\end{equation}  \\noindent where $\\Sigma_{HI}$ is the atomic hydrogen surface density, $\\Sigma_c$ is the critical surface density above which the disk is unstable to fragmentation, $N_{HI}$ is the \\HI~column density in units of 10$^{21}\\,$cm$^{-2}$, $v_{rot}$ is the rotation velocity in \\kms, and $r_{pc}$ is the orbital radius in parsecs. In the case of B2352+495 we obtain $\\Sigma_{HI}/\\Sigma_{c} \\sim 10$, thus, the disk is expected to be unstable to fragmentation. We note that reducing the assumed spin temperature by an order of magnitude would not significantly modify this conclusion, i.e., $\\Sigma_{HI}/\\Sigma_{c}$ would still be $\\sim 1$. This simple stability test indicates that the \\HI~cannot be distributed in a homogeneous disk. Indeed, the opacity variations from channel to channel (Figure~2) suggests that the \\HI~gas is not homogeneously distributed. We therefore concur with the standard schematic of clumpy atomic gas around AGNs (e.g., Peck et al.$\\,$1999).      \\subsection {Comparison with other High Angular Resolution Studies}  In this section we discuss the properties of \\HI~absorption in B2352+495 with respect to other \\HI~extragalactic absorbers that have been studied at very high angular resolution (i.e., interferometric observations with baselines $> 100\\,$km; Table~1). As is clear from Table~1 and the discussion in $\\S 1.1$, high angular resolution observations of \\HI~in CSOs reveal that the absorption does not trace a specific kind of cloud, instead, the absorbing clouds appear to be in quite different locations throughout the host galaxy, from infalling high velocity clouds and giant molecular clouds at kiloparsec scales, to torus and disks closer ($\\la 100\\,$pc) to the nucleus.  Only two of the five CSOs imaged at high angular resolution exhibit \\HI~absorption toward both radio lobes (Table~1; B1934$-$638 has not be imaged at high angular resolution, see $\\S$1.1). In general (four out of six, Table~1), CSOs studied at high angular resolution are characterized by a broad ($\\Delta V \\ga$ 100\\kms) \\HI~absorption component, and tend to be multi-peaked, with narrow ($\\Delta V < $ 100\\kms) lines detected in five of six cases.  High angular resolution observations of \\HI~toward non-CSO radio galaxies (Table~1) show a similar trend as the CSOs counterparts, i.e., the \\HI~absorption appears to trace a variety of environments, from disks directly associated with the AGN to clouds at kiloparsec-scale distances from the active nucleus. An apparent difference is the occurrence of \\HI~clouds interacting with radio jets. In the case of several non-CSO with double jet/lobe radio emission, the \\HI~appears to trace regions of interaction between atomic gas and radio jets, whereas such an interpretation has not been preferred in any of the CSOs imaged in \\HI~at high angular resolution. This trend does not appear to be due to bias of the researchers but, in most cases, seems due to real differences in the samples. For instance, the \\HI~absorption associated with tori/disks in CSOs is found very close ($\\la 100\\,$pc) to the radio core, close to the systemic velocity, and/or has evidence of rotation. Conversely, the \\HI~in non-CSO double jet/lobe sources with possible jet/cloud iteration is found at larger distances from the core ($\\ga500\\,$pc), shows no evidence of rotation, and/or has significantly different velocity from systemic. Although a more robust statistical sample is needed to achieve strong conclusions, the apparent absence of \\HI~interaction with the radio jets supports the current view that CSOs are young radio galaxies instead of being confined by dense circumnuclear material as proposed in the frustration scenario.  Arguments for the youth scenario are mostly based on the overall radio continuum properties of the sources, power budget estimates of the jets, expansion velocities, and properties of the nuclear molecular/atomic gas. CSOs lack strong extended radio emission related to energy deposited by the jets over long periods, and VLBI studies reveal proper motions that imply short dynamic ages (of the order of $10^3\\,$yr, e.g., Gugliucci et al. 2005). In addition, if the advance of the radio lobes would be truly frustrated, then for them to be in ram-pressure equilibrium the density of the surrounding medium should be above $10^2\\,$cm$^{-3}$. However, if the high density material is confined to a disk/torus then the frustration scenario is not viable. B2352+495 has evidence of being a young source based on all these arguments: it has no detected extended radio emission, the expansion velocity implies a dynamic age of $\\sim 1200\\,$yr (Taylor et al.$\\,$2000), and, as discussed above, the high density atomic gas appears to be distributed along a disk or torus which does not interact with the radio jet. Hence, the evidence favors the youth interpretation in this case."
    },
    "0910/0910.5729_arXiv.txt": {
        "abstract": "Using high resolution hydrodynamical simulations, we explore the spin evolution of massive dual black holes orbiting inside a circumnuclear disc, relic of a gas-rich galaxy merger.  The black holes spiral inwards from initially eccentric co- or counter-rotating coplanar orbits relative to the disc's rotation, and accrete gas that is carrying a net angular momentum.  As the black hole mass grows, its spin changes in strength and direction due to its gravito-magnetic coupling with the small-scale accretion disc.  We find that the black hole spins loose memory of their initial orientation, as accretion torques suffice to align the spins with the angular momentum of their orbit on a short timescale ($\\simlt 1-2$ Myr). A residual off-set in the spin direction relative to the orbital angular momentum remains, at the level of $\\simlt 10^o$ for the case of a cold disc, and $\\simlt 30^o$ for a warmer disc. Alignment in a cooler disc is more effective due to the higher coherence of the accretion flow near each black hole that reflects the large-scale coherence of the disc's rotation. If the massive black holes coalesce preserving the spin directions set after formation of a Keplerian binary, the relic black hole resulting from their coalescence receives a relatively small gravitational recoil.  The distribution of recoil velocities inferred from a simulated sample of massive black hole binaries has median $ \\simlt 70 \\kms$, much smaller than the median resulting from an isotropic distribution of spins. ",
        "introduction": "The massive black holes (MBHs) that we observe today in local spheroids (Ferrarese \\& Ford 2005, and references therein) are expected to have grown through a series of major accretion episodes in symbiosis with the growth of their host galaxies. Gas-rich major mergers may be at the heart of this joint evolution as they may explain the morphology of the hosts and at the same time account for the fueling of the underlying MBHs (e.g. Di Matteo, Springel \\& Hernquist 2005, and references therein).  In the currently favored cold dark matter hierarchical cosmologies, galaxy mergers play indeed a key role in growing galaxies to their present sizes and the coalescence of MBHs in binaries is therefore expected to be relatively common (Menou, Haiman \\& Narayanan 2001; Volonteri, Haardt \\& Madau 2003).  Following a galaxy major merger, the pair of MBHs first interacts with stars (Begelman, Blandford \\& Rees 1980; Milosavljevic \\& Merritt 2001) and gas (Mayer et al. 2007). The pair loses orbital angular momentum, and it is expected to harden progressively down to subparsec scales where emission of gravitational radiation drives the MBH inspiral down to coalescence. Depending on the properties of the coalescing binary, the pattern of the gravitational wave emission can be anisotropic, resulting in a non zero recoil velocity (a ``kick'') of the MBH remnant (Redmount \\& Rees 1989).  Several attempts to compute analytically the strength of the kick have been undertaken (Peres 1962; Bekenstein 1973; Fitchett 1983; Fitchett \\& Detweiler 1984; Redmount \\& Rees 1989; Wiseman 1992; Favata, Hughes \\& Holz 2004; Blanchet, Qusailah \\& Will 2005; Damour \\& Gopakumar 2006; Schnittman \\& Buonanno 2007.   Recent numerical simulations of the coalescence of spinning MBHs in full general relativity have been able to calculate explicitly kick velocities for a series of different binary configurations. It is found that three parameters influence the magnitude of the gravitational recoil of the relic MBH: the binary mass ratio, the spins, and the mutual orientation of the spins with respect to the orbital angular momentum. The recoil is largest, up to $4000\\, \\rm{km\\,s^{-1}}$, for nearly equal mass MBHs with large spins, when the spin vectors have opposite directions and are in the orbital plane (Campanelli et al. 2007).  By contrast, recoils of $\\lsim 200 \\kms$ are imparted to the MBH remnant if the spins of the progenitors prior coalescence are orthogonal to their orbital plane.  Purely general relativistic (GR) effects (i.e., spin-orbit and spin-spin interactions) may produce low recoil configurations, when the MBH pairing is driven by gravitational wave emission. However, those GR effects depend strongly on the initial relative orientation of the MBH spins, and can result in low recoil configurations only for a small region of parameter space (Schnittman 2004; Herrmann et al. 2009; Lousto et al. 2009).   In gas-rich mergers between galaxies of comparable mass (i.e. major mergers), close binary MBHs form under the action of dynamical friction against the gaseous and stellar background (Mayer et al. 2007; Callegari et al. 2009; see Colpi \\& Dotti 2009 for a review). During their inspiral MBHs are surrounded by a dense cocoon of gas that drives their dynamical decay and provides fuel for the feeding of the holes (Dotti et al. 2007, 2009).  Since matter carries angular momentum also the spin vector can change during the accretion process. The details of the dynamics may have a profound influence on the mass and spin evolution of the two MBHs, and thus on the recoil velocity of the MBH resulting from their coalescence, and this is matter of our concern in this paper.  The spin evolution during the MBH inspiral in a gas rich merger remnant has many implications .  Spins, prior to coalescence, influence the extent of the gravitational recoil, and so the retention of the relic MBH inside its host galaxy.  Accordingly, the spin distribution of the coalescing binaries, is critical as it determines the frequency of MBH retention in the host halo (Volonteri \\& Rees 2006; Volonteri 2007; Volonteri, Haardt \\& Gultekin 2008; Gultekin et al. in preparation).  The magnitude of the MBH spins and their orientation relative to the orbit during mergers is also critical in shaping the stellar density profiles in ellipticals, as the kicked MBH moving on a return orbit can deposit its excess kinetic energy into the stellar background, causing the formation of stellar core (Boylan--Kolchin, Ma \\& Quataert 2004; Gualandris \\& Merritt 2008). A recoiling MBH can have an observational signature when moving across the host galaxy, creating an X-ray tail in the perturbed hot gas (Devecchi et al. 2009), an off-set active nucleus (Loeb 2007; Volonteri \\& Madau 2008), shocking the inner rim of the accretion disc (Lippai, Frei \\& Haiman 2008; Schnittman \\& Krolik 2008), or dragging a stellar cusp with peculiarly high velocity dispersion (Merritt, Schnittman \\& Komossa 2009). Furthermore, spin orientations have important implications for gravitational wave astronomy, and for using gravitational wave measurements to constrain the formation history of MBHs (Vecchio 2004; Lang \\& Hughes 2006; Berti \\& Volonteri 2008; Arun et al. 2009a, 2009b).   Bogdanovic, Reynolds \\& Miller (2007) proposed a physical process that could {\\it align} the MBH spins with the orbital angular momentum of the binary, thus leading to slow recoils for the MBH remnant. The key process for alignment is the presence of a coherent gas inflow.  They speculate that accreting gas exerts gravito-magnetic torques that suffice to align the spins of both the MBHs with the angular momentum of the large-scale gas flow in which the orbit is embedded.   Spin--disc alignment due to gravito--magnetic coupling has been studied by a number of authors in the case of isolated MBHs surrounded by their own discs (Bardeen \\& Petterson 1975; Natarajan \\& Pringle 1998; Scheuer \\& Feiler 1996, Martin, Pringle \\& Tout 2007; Perego et al. 2009).  Here we attempt to explore for the first time spin-disc alignment around MBH binaries. We expect low recoils when the spin--disc coupling is strong, i.e. when: \\\\$\\bullet$ The two MBHs accrete $\\gsim 1 \\%$ of their initial mass before the coalescence (Natarajan \\& Pringle 1998; Natarajan \\& Armitage 1999; Volonteri, Sikora \\& Lasota 2007; Perego et al. 2009); \\\\$\\bullet$ The accreted gas carries angular momentum in a preferred direction, flowing onto each MBH along a preferential plane determined by the distribution of angular momentum of the gas in the environment of the MBH.\\\\  The first requirement can be fulfilled if the MBHs pair inside a dense, massive gaseous nuclear disc (Dotti et al. 2009), such as that predicted to form in remnants of gas-rich major mergers by Mayer et al. (2007).  To constrain the second requirement, we analyse a set of 3D Smooth Particle Hydrodynamics (SPH) simulations already discussed in Dotti et al. (2009). The high resolution of these simulations enables us to resolve the gravitational sphere of influence of each MBH during their inspiral inside the circumnuclear disc, and to map the distribution of angular momentum of the SPH particles in the MBH vicinity. MBHs are modeled as {\\it sink} particles that can accrete gas particles, allowing us to constrain the amount of mass accreted onto each MBH, and the orientation of the MBH spins relative to the angular momentum of the accreted gas.  The paper is organized as follows: in Section~2 we focus on the SPH simulations, and describe the semi-analytical algorithm that evolves the MBH spins; in Section~3 we illustrate our results on the MBH alignment and our prediction for the recoil velocity of the MBH remnant; in Section~4 we present our conclusions. ",
        "conclusions": "In this paper, we traced for the first time the evolution of the spin vectors of MBHs orbiting inside a massive circumnuclear gas disc. Our SPH simulations have sufficiently high resolution to probe the hydrodynamics of fluid particles and the accretion physics near the gravitational sphere of influence of the MBHs. An ad-hoc algorithm designed for tracking the gravo--magnetic coupling between the MBH spin and the small-scale accretion disc is then implemented in the code. We find that:  \\noindent $\\bullet$ When evolving a in dense, rotationally supported, structure such as a circumnuclear disc, MBHs in a pair align their spins ($\\mathbf J_{\\rm BH_{1,2}}$) to the pair orbital angular momentum ($\\mathbf L_{\\rm pair}$) well before the two MBHs bind in a Keplerian binary, and independently of the MBHs initial orbital parameters.  For a run with $M_2$ initially on a retrograde orbit, the spin of the secondary aligns efficiently only after the ``orbital angular momentum flip''.  \\noindent $\\bullet$ The average angle between $\\mathbf J_{\\rm BH_{1,2}}$ and $\\mathbf L_{\\rm pair}$ after the binary formation ($\\theta_{\\rm f}$) depends on the thermodynamics of the massive circumnuclear discs. $\\theta_{\\rm f}$ is lower if the MBHs are embedded in colder discs (with a polytropic index $\\gamma=7/5$), with respect to hotter discs ($\\gamma=5/3$);  \\noindent $\\bullet$ After the formation of a binary, the two MBHs accrete gas with the same dynamical and thermodynamical properties. As a consequence, even the angle between the two small projections of $\\mathbf J_{\\rm BH_{1,2}}$ in the orbital plane decreases. This further reduces the recoil velocity of the MBH remnant. The degree of alignment between the two spins and between each spin and $\\mathbf L_{\\rm pair}$ is preserved (or even increased) by spin--spin and spin--orbit interactions until the plunge phase (Schnittman 2004; Herrmann et al.  2009);   \\noindent $\\bullet$ Due to the efficient alignment between $\\mathbf J_{\\rm BH_{1,2}}$ and $\\mathbf L_{\\rm pair}$, the expected recoil velocities ($V_{\\rm kick}$) at the MBH coalescence is, on average, one order of magnitude lower than those expected for randomly oriented MBH spins.  The thermodynamical properties of the environment affect the degree of alignment and, as a consequence, the expected recoil velocities. $V_{\\rm kick}$ is lower (by a factor of 1.5--2.2) for lower values of $\\gamma$;   \\noindent $\\bullet$  Assuming the same distribution of $\\theta_{\\rm f}$, the recoil velocity distributions obtained using different prescriptions are very similar. The three fitting formulae used predict the same mean and median recoil velocities ($\\lsim 70 \\kms$) within a factor of $\\approx 1 - 1.3$, with less than few percent of realizations having $V_{\\rm kick}>400 \\kms$.   The distributions of recoil velocities that we find have important consequences for retention of MBHs in galactic nuclei.  When MBH binaries form and evolve in gas--rich major mergers, we predict the recoil velocity to be, on average, well below the escape speed from low-redshift galaxies. Indeed, because of the extreme efficiency of the spin alignment process, the recoil velocities are likely unimportant even for high-z proto-galactic building blocks.  Volonteri \\& Rees (2006) and Volonteri (2007) discussed how strong recoils can affect the early growth of MBHs at the highest redshifts.  In a forthcoming paper we will update our calculations and determine the impact of low recoils on the growth of MBHs in galaxies.   Our simple treatment of thermodynamics and the absence of any prescription for star--formation and supernovae feedback in our simulations could, in principle, overestimate the degree of coherency in the gas flows accreting onto the MBHs. Furthermore, our finite resolution prevent us to study the fragmentation of the accretion discs forming around the MBHs, that could result in a sequence of short and randomly oriented accretion events (King \\& Pringle 2006). We will investigate interaction between star formation in the circumnuclear disc and the properties of the accreting gas in a forthcoming study."
    },
    "0910/0910.3669_arXiv.txt": {
        "abstract": "Using a microlensing analysis of 11-years of OGLE V-band photometry of the four image gravitational lens Q2237+0305, we measure the inclination $i$ of the accretion disk to be $\\cos i > 0.66$ at 68\\% confidence. Very edge on ($\\cos i < 0.39$) solutions are ruled out at 95\\% confidence. We measure the V-band radius of the accretion disk, defined by the radius where the temperature matches the monitoring band photon emission, to be $\\RV = 5.8^{+3.8}_{-2.3}\\times 10^{15}\\cm$ assuming a simple thin disk model and including the uncertainties in its inclination. The projected radiating area of the disk remains too large to be consistent with the observed flux for a $T\\propto R^{-3/4}$ thin disk temperature profile. There is no strong correlation between the direction of motion (peculiar velocity) of the lens galaxy and the orientation of the disk. ",
        "introduction": "\\label{sec:intro}  In the AGN unification model \\citep[e.g.][]{Antonucci93,Urry95}, orientation differences among intrinsically similar objects are thought to account for many of the different observational properties of AGN. In particular, a dusty ``torus'' may frequently obscure the central engine from direct observation when viewed edge-on. The bright, Type 1 broad line quasars are thought to be viewed mostly face-on, so that the accretion disk is not obscured by material in the equatorial plane, and Type 2 narrow line quasars are viewed closer to edge-on through the obscuring material. Thus, if we could resolve the disk of a Type 1 quasar, we would expect it to be closer to face-on than edge-on. Unfortunately, familiar methods cannot resolve accretion disks, so we have only indirect measures of AGN disk orientation. For example, the projected axes of radio jets \\citep{Blandford79} and ionization cones \\citep{Elvis00} both support this picture. There are, however, no actual measurements of disk orientation.  While quasar accretion disks are too small to be resolved by direct imaging, gravitational microlensing provides a natural telescope to study the structure of quasar accretion disks and the properties of cosmologically distant lens galaxies where we see multiple images of background quasars \\citep[see][]{Wambsganss06}. In addition to the mean potential of the lens galaxy, each image is also magnified by the microlensing effects of the nearby stars. Since the observer, the lens galaxy and its stars, and the quasar are all moving, microlensing is observed as uncorrelated time variability in each of the quasar images. The amplitudes of these variations depend on the structure of the accretion disk and the properties of the lens galaxy.  Quasar microlensing is most sensitive to the projected area of the accretion disk relative to the source plane Einstein radius, \\begin{eqnarray} \\rE &=& D_{\\rm OS} \\sqrt{\\frac{4G\\left<M\\right>}{c^2}\\frac{D_{\\rm LS}}{D_{\\rm OL}D_{\\rm OS}}} \\nonumber \\\\ &=& 1.8\\times 10^{17} \\left(\\frac{\\left<M\\right>}{M_\\Sun}\\right)^{1/2}\\cm, \\end{eqnarray} where $G$ is the gravitational constant, $c$ is the speed of light, $\\left<M\\right>$ is the mean stellar mass of the stars, $D_{\\rm LS}, D_{\\rm OL}$, and $D_{\\rm OS}$ are the angular diameter distances between the lens-source, observer-lens, and observer-source respectively, and we have used the lens and source redshifts for Q2237+0305 \\citep[$z_{\\rm l}=0.0394\\,,z_{\\rm s}=1.685$,][Q2237 hereafter]{Huchra85}. The smaller the accretion disk, the higher the variability amplitude from microlensing. In general, the emission profile is difficult to determine, as models having similar half light radii show similar microlensing variability \\citep{Mortonson05}. There has been little examination of other structural parameters of disks, except for \\citet{Congdon07}, who demonstrated in simulations that the microlensing signal is sensitive to the ellipticity and orientation of the accretion disk. If detectable in practice, measuring the apparent ellipticities of accretion disks provides an important test of AGN unification models and opens the possibility of examining the complex effects of relativity on the apparent surface brightness of the disk \\citep[e.g.][]{Hubeny01}.  Measurements of quasar disk sizes using microlensing are now common. Recent efforts have studied individual sizes \\citep[e.g.][]{Morgan08a}, the relationships between size and wavelength \\citep{Anguita08,Bate08,Eigenbrod08b,Poindexter08,Floyd09,Mosquera09}, size and black hole mass \\citep{Morgan10}, and the sizes of thermal and non-thermal emission regions \\citep{Pooley07,Morgan08b,Chartas09,Dai09}. All these studies used circular accretion disks and static magnification patterns that neglect the random motion of stars in the lens galaxy.  \\begin{figure*} \\plotone{tracks2} \\caption{Example of a trial source trajectory (dark line segments) superposed on instantaneous point-source magnification patterns for $\\mmass = 0.3$. Darker shades indicate higher magnification. An HST H-band image in the center labels the images and the corresponding magnification patterns. Each pattern is rotated to have the correct orientation relative to the lens. This particular LC2 trial has an effective lens-plane velocity of $\\sim 600\\kms$ Northeast. The ``bow tie'' (top) exhibits the two peaks in the major-axis position angle distribution together with the shaded region it approximately represents the 68\\% confidence region. The solid disk (right) has a radius of $10^{17}\\cm$.} \\label{fig:patterns} \\end{figure*}  Determining the shape and orientation of a disk depends on the correlations between the anisotropic structure of the disk and the anisotropic structure in the magnification patterns created by the shear (tidal gravity) local to each image (see Figure \\ref{fig:patterns}, \\citealt{Congdon07}). Existing microlensing studies cannot safely explore these issues because they neglect the motions of the stars in the lens galaxy and use ``static'' magnification patterns. Since the stellar velocity dispersions of lens galaxies are comparable to the peculiar velocities of galaxies, the patterns change on the same time scale as the source traverses them. Ignoring these stellar motions will overestimate the coherence of the magnification patterns \\citep[see][]{Wyithe00,Kochanek07} and likely render estimates of disk shapes unrealistic. With a few exceptions that do not focus on disk structure (see Paper I and \\citet{Wyithe99}), analyses of microlensing data have used static magnification patterns because of the computational challenges. In Paper I \\citep{Poindexter09}, we solved these computational problems and examined the peculiar velocity of the lens galaxy of Q2237 and the mean mass of its stars. In this paper, we measure the size, inclination, and position angle of the accretion disk in Q2237. In \\S\\ref{sec:data} we describe the data set used for our microlensing analysis, our disk model, and outline our overall approach. Our results are presented in \\S\\ref{sec:results}, and a discussion follows in \\S\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  By including random stellar motions in our microlensing analysis of Q2237 we find evidence that the accretion disk of Q2237 is viewed face-on with $\\cos i > 0.63$. This lends support to the popular AGN unification model where we expect Type 1 quasars like Q2237 to be viewed nearly face-on. Including 20\\% contamination from broad line emission on larger spatial scales \\citep[estimated from spectra][]{Eigenbrod08a} results in a stronger case for face-on solutions. Modeling stellar motions and the inclination and position angle parameters further reduces the systematic uncertainties of our measurements of the disk radius compared to earlier studies by including a broader range of physical uncertainties. As we found in Paper I for the lens velocity and mean stellar mass, the results of the separate analyses of LC1 and LC2, the first and second temporal halves of the OGLE light curves, produce consistent results for every parameter we considered. While we have used a relatively simple model for the accretion disk, these results demonstrate that disk shapes can be measured with quasar microlensing, as suggested by \\citet{Congdon07}. As data sets and computing power improve, it will be natural to try fitting more subtle disk features such as the asymmetries from relativistic effects using relativistic models such as \\citet{Hubeny01}.  The exceptionally long OGLE light curve and the fast microlensing timescales of Q2337 made it a natural first choice for including the random stellar motions and studying the shape of the disk. Our expanded method is similar in computational cost to our previous efforts with static patterns, so there is no reason not to use it generally. The stellar motions clearly aid in reducing uncertainties in the mean mass (see Paper I) and allow us to correctly make inclination corrections. It can also easily be extended to take advantage of multi-wavelength data to try to constrain the temperature profile. However, as in earlier microlensing studies \\citep{Pooley07, Morgan10, Dai09}, we cannot reconcile the basic temperature profile of a thin disk, the microlensing size estimate, and the observed optical flux.  A disk with a $T\\propto R^{-3/4}$ temperature profile normalized by the microlensing size estimate should be brighter than observed.  This can be solved by altering the temperature profile \\citep[see][]{Poindexter08,Morgan10}. For example, reducing the slope from $T\\propto R^{-3/4}$ to $T \\propto R^{-1/2}$ would increase the flux size by a factor of $2.2$ relative to the half light radius. This would be mildly inconsistent ($1.5\\sigma$) with the slope estimate of $-0.83 \\pm 0.21$ by \\citet{Eigenbrod08b}. However such changes also call into question the basic structure of the thin disk model. The other simple possibility is to reduce the emissivity of the disk to be well below that of a blackbody (by the ratio of the flux/microlensing sizes squared), but this seems less physically plausible than change in the temperature structure."
    },
    "0910/0910.1346_arXiv.txt": {
        "abstract": "Radiative transfer in planetary atmospheres is usually treated in the static limit, i.e., neglecting atmospheric motions. We argue that hot Jupiter atmospheres, with possibly fast (sonic) wind speeds, may require a more strongly coupled treatment, formally in the regime of radiation-hydrodynamics. To lowest order in $v/c$, relativistic Doppler shifts distort line profiles along optical paths with finite wind velocity gradients. This leads to flow-dependent deviations in the effective emission and absorption properties of the atmospheric medium. Evaluating the overall impact of these distortions on the radiative structure of a dynamic atmosphere is non-trivial. We present transmissivity and systematic equivalent width excess calculations which suggest possibly important consequences for radiation transport in hot Jupiter atmospheres. If winds are fast and bulk Doppler shifts are indeed important for the global radiative balance, accurate modeling and reliable data interpretation for hot Jupiter atmospheres may prove challenging: it would involve anisotropic and dynamic radiative transfer in a coupled radiation-hydrodynamical flow. On the bright side, it would also imply that the emergent properties of hot Jupiter atmospheres are more direct tracers of their atmospheric flows than is the case for Solar System planets. Radiation-hydrodynamics may also influence radiative transfer in other classes of hot exoplanetary atmospheres with fast winds. ",
        "introduction": "\\label{S:intro}   In recent years, it has become possible to remotely observe the atmospheres of some exoplanets found in close-in orbits around nearby stars. Many of these observations, which include eclipse, phase curve, and transit photometric or spectroscopic measurements, have focused on the specific class of exoplanets known as hot Jupiters \\citep{Charbonneau2002,Charbonneau2005,Deming2005,Charbonneau2008,Knutson2008,Knutson2009b,Harrington2006,Harrington2007,Agol2009,Cowan2007,Knutson2007,Knutson2009a,Tinetti2007,Pont2008,Sing2008,Swain2008a,Swain2008b,Grillmair2007,Richardson2007}. These gaseous giant planets have high atmospheric temperatures, from the close proximity to their parent stars, slow spin rotation rates, as inferred from their expected state of tidal synchronization, and they are thus subject to an unusual steady pattern of hemispheric insolation, with permanent day and night sides \\citep[e.g.,][]{Seager2005,SMC08,SCM10}.   A variety of diagnostics about physical conditions in the atmospheres of a few specific hot Jupiters have been presented in the literature, from the presence of high altitude haze \\citep[e.g.,][]{Pont2008} to the existence of vertical temperature inversions \\citep[\"stratospheres\", e.g.,][]{Burrows2008,Fortney2008}, as well as constraints on the chemical abundances of radiatively active species \\citep[e.g.,][]{Charbonneau2002,Tinetti2007,Sing2008,Swain2008b}. In most cases, these interpretations rely on one-dimensional, steady-state radiative transfer models describing in great detail the atmospheric structure that is expected at radiative and chemical equilibrium. On the other hand, these models usually consider only globally averaged atmospheric properties and they often ignore the role of horizontal heat transport by advection on the energetic balance of the modeled atmospheres. Such shortcomings of radiative transfer models could eventually limit our ability to reliably infer the physical conditions present in remotely observed exoplanetary atmospheres.   Planetary atmospheres which satisfy global radiative equilibrium are generally not in local radiative equilibrium, that is along a specific vertical column, because of finite contributions from horizontal heat fluxes. The role of atmospheric circulation in shaping the structure and properties of hot Jupiter atmospheres has been recognized as an increasingly important ingredient over the last few years \\citep[e.g.,][]{Fortney2006} and several groups have been developing multi-dimensional atmospheric models to address this issue more explicitly \\citep{SG02,Cho2003,Cho2008,Burkert2005,CS05,CS06,Langton2007,DobbsDixon2008,S08,S09,MR08,RauMen09}.   Even in a three-dimensional, time-dependent atmospheric model of the type known as \"General Circulation Model\" (GCM), radiative transfer is traditionally treated in the static approximation. In this limit, the coupling between the atmospheric flow and the radiative heat transport enters only via the energy equation for the moving atmospheric fluid, so that this coupling is purely thermodynamic in nature. The radiative transfer component of the problem, which is needed to continuously evaluate the net diabatic heating or cooling rate of the gas in motion, is solved independently of the atmospheric motions themselves, as if the atmosphere were actually static.  Here, we argue that this level of thermodynamic coupling, which is a standard approximation for Solar System planetary atmospheres, may be insufficient to describe accurately the nature of energy transport in hot Jupiter atmospheres. A more strongly coupled treatment known as radiation-hydrodynamics, in which flow velocities enter the radiation transport problem explicitly, may be required to describe radiative transfer in hot Jupiter atmospheres with fast, sonic or transonic wind speeds.   The remainder of this paper is organized as follows. In \\S\\ref{S:dynatm}, we outline the general formalism for radiation transport in a dynamic atmosphere. We motivate the relevance of this radiation-hydrodynamics regime for hot Jupiter atmospheres with fast winds in \\S\\ref{S:HotJups}. We discuss the potential magnitude of deviations from the static limit for the global radiative balance of hot Jupiter atmospheres in \\S\\ref{S:EW} and we conclude in \\S\\ref{S:impl}. ",
        "conclusions": "\\label{S:impl}   Our main result has been to establish the possibility that bulk Doppler shifts from fast, sonic or transonic wind speeds in hot Jupiter atmospheres could affect radiation transport and thus in principle influence the overall radiative structure of these atmospheres. We have quantified the magnitude of these effects with simplified transmissivity models in a dynamic atmosphere, by assuming that linear bulk velocity gradients are present along otherwise homogeneous optical paths and by adopting Voigt line profile parameters appropriate for the radiatively forced regions of hot Jupiter atmospheres. We have found that bulk Doppler shifts systematically increase the equivalent widths of isolated radiative lines, relative to the equivalent widths in a static atmosphere, with excesses easily reaching several tens of percents for relevant model parameters.   Flux errors below the few percent level may well be acceptable in atmospheric models currently used to interpret the increasingly rich set of observational data on hot Jupiter atmospheres since other important atmospheric unknowns also contribute to modeling errors (e.g., the exact compositional profiles of radiatively active constituents). Flux errors at the level of several tens of percents or more may be too large to be ignored, however, especially if they originate from systematic excesses in the equivalent widths of important radiative lines. Given the link between line equivalent widths and radiative intensities discussed in \\S\\ref{sec:EW} in the context of the narrow-band formalism, we therefore suggest that equivalent width excesses of the magnitude found in several of our idealized models could have a significant impact on radiative fluxes and the overall radiative energy balance of hot Jupiter atmospheres.   However, we must also caution that, beyond these preliminary arguments, it is not possible to settle this question without a considerably more detailed calculation than the one presented here. Indeed, it is the combined effect of bulk Doppler shifts on a vast number of radiative lines, all with varying degrees of optical thickness, weighted by the various Planck functions of the atmospheric layers under consideration (thus depending on the atmospheric temperature profile itself), along a variety of inhomogeneous optical paths, which ultimately determines by how much the overall radiative balance of the dynamic atmosphere is affected. By contrast, our simple transmissivity and equivalent width models were focused on single isolated lines with purely linear velocity gradients along arbitrary, homogeneous optical paths.  Furthermore, contributions from continuum opacity sources in the atmosphere, for instance in cloudy regions, could reduce the impact of bulk Doppler shifts on the atmospheric radiation transport, since small shifts mostly affect narrow radiative lines. Various simplifying assumptions made in our models would thus be invalidated in a more realistic radiative transport calculation. Interestingly, the full radiation transport problem in a dynamic atmosphere may be amenable to practical numerical solutions in the future \\citep[e.g.,][]{Knop2009}.  It is also worth emphasizing that our discussion has largely focused on the long-wavelength (thermal) component of the atmospheric radiation spectrum, for which the diffuse approximation and the use of an isotropic (Planck) source function are justified. At first, one may be tempted to neglect the effects of bulk Doppler shifts from horizontal winds in the treatment of short-wavelength atmospheric radiation since the corresponding beam-like insolation is vertical in first approximation. This could be misleading, however. When scattering is important, as is usually the case, it generates optical paths slanted enough that they could become much more susceptible to the bulk Doppler shifts caused by horizontal winds. Furthermore, there is a significant fraction of the atmosphere that is located far enough away from the substellar point to be subject to partially slanted irradiation. In the case of close-in planets like hot Jupiters, the finite size of the stellar disk could also contribute to slanted irradiation and thus an increased sensitivity to Doppler shifts from horizontal winds. Let us emphasize that radiation transport in a dynamic atmosphere with an anisotropic source function, as would result from multiple scattering of the beam-like insolation in the short-wavelength portion of the atmospheric radiation spectrum, would lead to a considerably more complex formulation of the radiative transfer problem than discussed here for the case of diffuse thermal radiation (e.g., Eq.~[\\ref{eq:two}]). For example, one wonders whether the anisotropic transport resulting from atmospheric bulk motions could result in a stronger polarization signal than has been estimated on the basis of static radiative transfer calculations \\citep[e.g.,][]{Seager2000}.   Interestingly, the various short-wavelength radiation transport effects mentioned above should be particularly pronounced in the context of transmission spectroscopic measurements, which specifically probe nearly horizontal optical paths at the atmospheric planetary limb during transits.  Essentially all transit spectroscopic diagnostics could thus be affected by bulk Doppler shifts from fast atmospheric winds. \\citet{Brown01} has presented a detailed discussion of this problem, including consequences of the vertical gradients of horizontal wind velocity. The possibly more important effects of bulk Doppler shifts along a specific optical path were omitted from these calculations, however, even though the author did comment on the expectation of increased line equivalent widths. It may thus prove important to carefully reassess various interpretations about the chemical composition of hot Jupiter atmospheres based on transit spectroscopic measurements with models which properly account for radiation transport in dynamic, rather than static, atmospheres.   Finally, let us conclude by mentioning a few possible extensions of this work to other classes of planetary atmospheres. It would seem that the standard static assumption for the treatment of radiation transport in Solar System planetary atmospheres is well justified. Indeed, for these much cooler atmospheres, at a representative pressure level of 1 bar, radiative lines are significantly more pressure-broadened in Solar System atmospheres than in hot Jupiter atmospheres (see scalings in \\S\\ref{sec:widths}). In that limit, and for wind speeds well below the sound speed, our models do indicate very small excesses in the equivalent width of radiative lines.\\footnote{Interestingly, Neptune's and Saturn's upper atmospheres and their near-sonic wind speeds \\citep[e.g.,][]{Suomi1991,Li08} may constitute a partial exception to this rule, despite cold atmospheric temperatures.} By contrast, the typically much hotter planets discovered by astronomers in recent years, which are often subject to unusually strong radiative forcing conditions, would seem to be more natural sites for the application of the radiation-hydrodynamical principles emphasized in the present work. Besides hot Jupiters, these principles could also find applications in the class of eccentric giant planets which experience transient atmospheric flash-heating during periastron passage \\citep[e.g.,][]{Langton2008,Laughlin09} and perhaps the emerging class of hot super-Earths, if the atmospheres of these planets are able to sustain wind velocities approaching or exceeding the sound speed in the presence of significant ground drag."
    },
    "0910/0910.4563_arXiv.txt": {
        "abstract": "A leading candidate for astrophysical dark matter (DM) is a massive particle with a mass in the range from 50 GeV to greater than 10 TeV and an interaction cross section on the weak scale.  The self-annihilation of such particles in astrophysical regions of high DM density can generate stable secondary particles including very high energy gamma rays with energies up to the DM particle mass. Dwarf spheroidal galaxies of the Local Group are attractive targets to search for the annihilation signature of DM due to their proximity and large DM content.  We report on gamma-ray observations taken with the Very Energetic Radiation Imaging Telescope Array System (VERITAS) of several dwarf galaxy targets as well as the globular cluster M5 and the local group galaxies M32 and M33.  We discuss the implications of these measurements for the parameter space of DM particle models ",
        "introduction": "\\label{sec:intro}  The existence of astrophysical non-baryonic dark matter (DM) has been established by its gravitational effect on galaxy rotation \\cite{rub80} and the velocity dispersions of objects from dwarf galaxies \\cite{wal07} through large galaxy clusters \\cite{zwi37}.  Additional evidence of DM existence comes from cosmic microwave background measurements \\cite{spe07} and gravitational lensing of galaxies by foreground galaxy clusters, for example, the evident separation of dark and baryonic matter in 1E0657-558, the ``bullet cluster'' \\cite{clo06}.  At present though, the existence of dark matter is solely inferred from its gravitational influence.  The particle nature of dark matter is yet to be revealed through direct detection in terrestrial dark matter searches, its production in particle accelerators, or indirect dark matter searches looking for evidence of a flux of known particles produced by annihilation of dark matter particle pairs.  Here we report on an indirect dark matter search carried out using VERITAS located at the Fred Lawrence Whipple Observatory in southern Arizona, USA, at an elevation of 1268m \\cite{wee02}.  Gamma rays can be produced either directly from the pair annhilation of weakly interacting massive particles (WIMPs) or as secondary decay products from the primary particles produced in WIMP annihilation. The former would produce a monoenergetic source of gamma rays equal to the WIMP mass and would constitute definitive evidence for particle dark matter.  The latter mechanism would produce a spectrum of gamma-ray energies with a cutoff at the WIMP mass.  For astrophysical targets, the WIMP annihilation rate and associated gamma-ray flux are highly uncertain due to theoretical uncertainties and limited observational constraints on the DM halo profile.  Therefore, VERITAS has surveyed a variety of possible sources in its dark matter search.  We present here results from observations of three dwarf spheroidal galaxies (dSph): Draco, Ursa Minor, and Willman 1; the globular cluster M5; and the local group galaxies M32 and M33.  The emphasis of the program has been to target dSphs due to their relative proximity and low expected background from known astrophysical gamma-ray sources.  VERITAS can detect and measure gamma rays in the $\\sim$100 GeV - 30 TeV energy range with an energy resolution of 15-20\\%, and angular resolution of $\\sim0.1^\\circ$ per event. Further technical description of VERITAS can be found in \\cite{acc08}. ",
        "conclusions": "\\label{sec:conclu}  We have carried out a search for very high energy gamma rays from three dwarf spheroidal galaxies: Draco, Ursa Minor, and Willman I; the globular cluster, M5; and the local galaxies M32 and M33 as part of an indirect dark matter search program on the VERITAS IACT array. No significant excess above background was observed from any target. We set upper limits on the flux and, for the dwarf spheroidal galaxies limits on the cross section times velocity, $\\langle\\sigma v \\rangle$, for neutralino pair annihilation as a function of neutralino mass.  The $\\langle\\sigma v \\rangle$ limits indicate that a substantial boost factor above smooth dark halo expectations would be required of MSSM-type models.  We will continue our program in the future with an emphasis on further dSph observations. Next generation IACT arrays now being planned such as the Advanced Gamma-ray Imaging System (AGIS) and the Cherenkov Telescope Array (CTA) will provide an order of magnitude increase in sensitivity over current arrays such as VERITAS, MAGIC II, and HESS and can be expected to constrain some models of both supersymmetric dark matter or Kaluza-Klein dark matter associated with models of universal extra dimensions even in the conservative smooth DM halo paradigm.  Along with observations by the Fermi Gamma-ray Space Telescope and new particle searches at the Large Hadron Collider, prospects for understanding the nature of dark matter over the next decade look to be promising.  This research is supported by grants from the US Department of Energy, the US National Science Foundation, and the Smithsonian Institution; by NSERC in Canada; by Science Foundation Ireland; and by STFC in the UK.  We acknowledge the excellent work of the technical support staff at the FLWO and the collaborating institutions in the construction and operation of the instrument.  \\begin{table*}[ht] \\begin{threeparttable} \\centering \\caption{Summary Results of Dwarf Galaxy Observations\\label{tbl:results}} \\begin{tabular}{|l|c|c|c|} \\hline Quantity                                 &  Draco  &  Ursa Minor  &  Willman I    \\\\ \\hline Exposure (hr)                            &  18.38  &    18.91     &    13.68      \\\\ Signal Region (events)                   &    305  &      250     &      326      \\\\ Total Background (events)                &   3667  &     3084     &     3602      \\\\ Number Backgrd. Regions                  &     11  &       11     &       11      \\\\ Significance\\tnote{a}                    &  -1.51  &    -1.77     &    -0.08      \\\\ 95\\% c.l. (counts)\\tnote{b}              &   18.8  &     15.6     &     36.7      \\\\ Effective Area ($\\mbox{m}^2$)            &  12518  &    16917     &    33413      \\\\ Energy Threshold (GeV)                   &    200  &      200     &      200      \\\\ Flux Limit 95\\% c.l. ($\\mbox{cm}^{-2} \\mbox{s}^{-1}$)  & $2.82 \\times 10^{-12}$  &  $1.35 \\times 10^{-12}$ &  $2.23 \\times 10^{-12}$     \\\\ \\hline \\end{tabular} \\begin{tablenotes} \\item[a] Li and Ma eqn.~17 method \\cite{li83} \\item[b] Rolke method \\cite{rol05} \\end{tablenotes} \\end{threeparttable} \\end{table*}   \\begin{table}[h] \\begin{threeparttable} \\caption{Parameters used for the astrophysical factor, $J$ calculation. \\label{tbl:J}} \\centering \\begin{tabular}{|l|c|c|c|} \\hline Quantity                                         &       Draco          &     Ursa Minor       &   Willman I       \\\\ \\hline $R_{dSph}$ (kpc)\\tnote{a}                        &        80            &       66             &     38            \\\\ $r_t$  (kpc)\\tnote{b}                            &         7            &        7             &      7            \\\\ $\\rho_s (\\mbox{M}_\\odot/\\mbox{kpc}^3)$\\tnote{c}  &   $4.5\\times 10^7$   &   $4.5\\times 10^7$   &   $4\\times10^8$   \\\\ $r_s$ (kpc)\\tnote{d}                             &      0.79            &     0.79             &    0.18           \\\\ $J$\\tnote{e}                                     &         4            &        7             &     22            \\\\ \\hline \\end{tabular} \\begin{tablenotes} \\item[a] Earth-dwarf galaxy distance \\item[b] The value of $J$ is negligibly changed for tidal radius, $r_t$, as low as 0.9 kpc. \\item[c] scale density \\item[d] scale radius \\item[e] $J$ is expressed as a dimensionless value normalized to the critical density squared times the Hubble radius, $3.832\\times10^{17} \\mbox{GeV}^2\\mbox{cm}^{-5}$. \\end{tablenotes} \\end{threeparttable} \\end{table}  \\begin{table}[h] \\begin{threeparttable} \\caption{Summary Results of Globular Cluster and Large Galaxy Observations\\label{tbl:other-results}} \\centering \\begin{tabular}{|l|c|c|c|} \\hline Quantity                                 &        M5            &      M32            &     M33           \\\\ \\hline Exposure (hr)                            &       15.0           &    11.29            &    11.83          \\\\ Signal Region (events)                   &         25           &      262            &      147          \\\\ Total Background (events)                &        251           &     2156            &      992          \\\\ Number Backgrd. Regions                  &         11           &        7            &        7          \\\\ Significance\\tnote{a}                    &       -0.3           &     0.59            &     0.41          \\\\ 95\\% c.l. (counts)\\tnote{b}              &       13.6           &     12.9            &     31.8          \\\\ Flux 95\\% c.l. (photons $\\mbox{s}^{-1}$) &  $2.5\\times10^{-4}$  & $3.2\\times10^{-4}$  &  $7.4\\times10^{-4}$  \\\\ \\hline \\end{tabular} \\begin{tablenotes} \\item[a] Li and Ma eqn.~17 method \\cite{li83} \\item[b] Rolke method \\cite{rol05} \\end{tablenotes} \\end{threeparttable} \\end{table}"
    },
    "0910/0910.4945.txt": {
        "abstract": "The measurement of redshifts for Gamma-Ray Bursts (GRBs) is an important issue for the study of the high redshift universe and cosmology. We are constructing a program to estimate the redshifts for GRBs from the original \\textit{Swift} light curves and spectra, aiming to get redshifts for the Swift bursts \\textit{without} spectroscopic or photometric redshifts. We derive the luminosity indicators from the light curves and spectra of each burst, including the lag time between low and high photon energy light curves, the variability of the light curve, the peak energy of the spectrum, the number of peaks in the light curve, and the minimum rise time of the peaks. These luminosity indicators can each be related directly to the luminosity, and we combine their independent luminosities into one weighted average. Then with our combined luminosity value, the observed burst peak brightness, and the concordance redshift-distance relation, we can derive the redshift for each burst. In this paper, we test the accuracy of our method on 107 bursts with known spectroscopic redshift. The reduced $\\chi^2$ of our best redshifts ($z_{best}$) compared with known spectroscopic redshifts ($z_{spec}$) is 0.86, and the average value of $log_{10}(z_{best}/z_{spec})$ is 0.01, with this indicating that our error bars are good and our estimates are not biased. The RMS scatter of $log_{10}(z_{best}/z_{spec})$ is 0.26, with a comparison of 0.30 for RMS of $log_{10}(z_{spec})$. We made a selection on bursts with relatively accurate redshift estimation. The RMS of $log(z_{best}/z_{spec})$ decreases to 0.19, and the RMS scatter of $log_{10}(z_{spec})$ for this subsample is 0.28. For \\textit{Swift} bursts measured over a relatively narrow energy band, the uncertainty in determining the peak energy is one of the main restrictions on our accuracy. Although the accuracy of our $z_{best}$ values are not as good as that of spectroscopic redshifts, it is very useful for demographic studies, as our sample is nearly complete and the redshifts do not have the severe selection effects associated with optical spectroscopy. ",
        "introduction": "The redshifts have long been an important issue for long-duration Gamma-Ray Bursts (GRBs). Since the first X-ray, optical, and radio counterparts were discovered in 1997 (Costa et al. 1997; van Paradijs et al. 1997; Frail et al. 1997), GRBs are confirmed to be in galaxies at cosmological distances. However, until now, only a small fraction of the bursts have their redshifts measured. Even \\textit{Swift}, over its first four years of operation, has only roughly 30$\\%$ of its bursts with spectroscopic/photometric redshifts. These redshifts may have complex selection biases as a function of redshift relating to the difficulty of getting redshifts for faint and distant bursts as well as the distribution of bands used for spectra and photometry. One illustration of the importance of the selection effects is that the average of redshift value for \\textit{Swift} GRBs have been dropped from $< z >$ = 2.8 (Jakobsson et al. 2006c) and to $< z >$ = 2.1 (Jakobsson 2008) while the average redshift value of earlier satellites is much lower yet (Bagoly et al. 2009). As the real redshift distribution of GRBs cannot be changed by that much for the past few years, the only reason for it would be changing selection effects in the spectroscopic observations.  These systematic biases will greatly affect various analyses. As long-duration GRBs come from the collapse of very massive fast-rotating stars with very short main-sequence lifetimes (Woosley $\\&$ Bloom 2006), their rate density will provide a measurement of the massive star formation rate of our Universe. If we only deal with bursts with known spectroscopic redshifts, the derived star formation rate may be biased by the spectroscopic-redshift selection effects. The same problem comes with demographic studies of the GRB luminosity function. If only bursts with spectroscopic redshifts are used in constructing the luminosity functions, then the selection biases may distort the derived luminosity function. A third important problem is the recognition of the highest redshift bursts (say, with $z>7$), against which the current spectroscopic methods may be heavily biased. \\textit{Swift} is expected to have $5-10\\%$ of its bursts with $z>7$ (Bromm $\\&$ Loeb 2006), but the spectroscopic methods have only identified one such burst GRB090423, and that just a few month ago (Tanvir et al. 2009).  A solution to these problems is to use the luminosity relations that work on luminosity indicators measured directly from the prompt gamma-ray emission. These luminosity relations are equations that connect a measurable property of the burst itself (the luminosity indicators) to the burst\u00a1\u00afs energy ($E_{\\gamma}$ or $E_{\\gamma,iso}$) or the peak luminosity (L). The observed fluence or peak flux can then be used (with the inverse-square law) to derive a distance to the burst and (for some fiducial cosmology) the burst redshift. The advantage of this procedure is that it applies to all long class GRBs, not just to the minority selected to have a measured spectroscopic redshift. The disadvantage is that the uncertainties on the derived redshifts are much larger than those of the spectroscopic redshifts, and so the method can be used primarily for statistical or demographic purposes. The situation is similar to the case where photometric redshifts of distant galaxies in the Hubble Deep Field do not have the accuracy of the spectroscopic redshifts, yet nevertheless these photometric redshifts can be done for all the galaxies and are the cornerstone of all statistical studies of the field.  The six luminosity relations we use in this article are: $\\tau_{lag}$ (spectral lag) - L relation (Norris, Marani, $\\&$ Bonnell 2000), V (Variability) - L relation(Fenimore $\\&$ Ramirez-Ruiz 2000), $E_{peak}$ (peak energy of the spectrum) - L relation (Schaefer 2003b), $E_{peak}$ - $E_{\\gamma}$ relation (also called Ghirlanda relation, Ghirlanda, Ghisellini, $\\&$ Lazzati (2004)), $\\tau_{RT}$ (minimum rise time) - L relation (Schaefer 2002), and $N_{peak}$ (number of peaks in the light curve) - L relation (Schaefer 2002). It was claimed by Butler et al. (2007) that the $E_{peak}$ - $E_{\\gamma,iso}$ relation (also called Amati relation, Amati et al. (2002)) differs somewhat between \\textit{Swift} and pre-\\textit{Swift} data, and hence the relation might be caused by the detection threshold effect of the instrument instead of the GRB itself. We tested this claim on the relations we are using, as shown in Section 3.  In this article, we will demonstrate a method to use these luminosity relations to estimate the GRBs redshifts. Here, we will only consider bursts with known spectroscopic redshifts, but our derivation of $z_{best}$ is based on the burst properties, and the effect of the known spectroscopic redshifts are negligible. This effect is tested later in the article. By applying our method to the bursts with accurate spectroscopic redshifts measured, we can compare our estimated redshifts with the measured spectroscopic redshifts. This will be the key test of our methods and the accuracy of our derived redshifts, as well as the applicability of our methods to bursts without spectroscopic redshift. With confidence in the reliability of the method, we can then apply it to all \\textit{Swift} bursts including those with no spectroscopic redshifts. Also, we will be able to apply this method to future bursts, and provide the community with a rapid notification about their redshifts. (Xiao $\\&$ Schaefer 2009) ",
        "conclusions": "In this paper, we developed a method to calculate the redshifts for long GRBs, using their light curves and spectra. We applied our method to bursts with known spectroscopic redshifts, detected by BATSE, HETE, Konus and \\textit{Swift}. By comparing our calculated redshifts with their spectroscopic redshifts, we are able to examine the accuracy of our method.  We compared each of the luminosity relations for pre-\\textit{Swift} and \\textit{Swift} bursts by making a F-test. With the F values close to unity, we have significant evidence against any claim that the relations are caused by the detection threshold effects or any other artificial effects of the instruments.  We compared our results with the spectroscopic redshifts. We find that our $z_{best}$ are not biased (with the average value of $log_{10}(z_{best}/z_{spec})$ equal to 0.01), and our reported $1-\\sigma$ error bars are good (with $\\chi_{red}^2 = 0.86$, and $70\\%$ of the $z_{spec}$ fall into the $1-\\sigma$ region of $z_{best}$). Our accuracy on the redshifts are not as accurate as those from spectroscopy, yet nevertheless with a reasonable accuracy for demographical and statistical studies, with the RMS of $log_{10}(z_{best}/z_{spec})$ is 0.26. The RMS value is about twice as what was found in Schaefer (2007). One of the reason is, in Schaefer (2007), the accuracy is calculated assuming a known $z_{spec}$, and in this paper, as we are treating the unknown-redshift case, so extra degrees of freedom has been brought in the calculation, which caused the accuracy to get worse by about a factor of 2. Another reason is that the large RMS is dominated by those faint and noisy burst, and for a subsample with $\\sigma_{z_{minus}/z_{best}}<0.5$ and $\\sigma_{z_{plus}/z_{best}}<1$, we get the RMS of $[log(z_{best}/z_{spec})]$ of 0.19, which is much smaller. As our $z_{best}$ are from the light curves, the spectra and the concordance cosmological model, it is independent of the spectroscopic redshift. As a result, our method can be applied to all long GRBs.  For \\textit{Swift} bursts, as we are measuring over a relatively narrow energy band (15 keV - 350 keV), the uncertainties in the calculation of peak energy $E_{peak}$ is large, and it becomes one of the main restrictions on our accuracy. With the launch of Fermi, we are hoping to get bursts with more accurately measured $E_{peak}$ values and light curves covering a broader energy band. We are expecting a substantial improvement in the accuracy of redshifts for Fermi bursts.  For the next step, we will apply our method to all \\textit{Swift} long GRBs, aiming to get a nearly-complete \\textit{Swift} GRBs redshift catalog. Such a catalog will inevitably be incomplete due to bursts with incomplete light curves and bursts too faint for their properties to be usefully measured. Our resulting redshifts will have an accuracy worse than those obtainable with optical spectroscopy, yet our accuracy will be good enough for various important statistical studies. As such, our catalog will be used for the demographic studies, without the detection threshold effect of the spectroscopic redshift measurements. We are also trying to deal with all the future bursts, and to provide a rapid notification of the redshift on the GCN circular to the community. We are hoping to find some possible high redshift (e.g. $z>7$) and possible low redshift (e.g. $z<0.3$) GRBs, which will be important in the observation of the high redshift universe and the GRB-SNe connection study."
    },
    "0910/0910.3944_arXiv.txt": {
        "abstract": "We report the most accurate X-ray position of the X-ray source in the giant globular cluster G1 in M31 by using the \\chandra\\ X-ray Observatory, {\\it Hubble Space Telescope (HST)}, and Canada-France-Hawaii Telescope (CFHT). G1 is clearly detected with \\chandra\\ and by cross-registering with \\hst\\ and CFHT images, we derive a $1\\sigma$ error radius of $0\\farcs15$, significantly smaller than the previous measurement by \\xmm. We conclude that the X-ray emission of G1 is likely to come from within the core radius of the cluster. We have considered a number of possibilities for the origin of the X-ray emission but can rule all but two scenarios out: it could be due to either accretion onto a central intermediate-mass black hole (IMBH), or an ordinary low-mass X-ray binary (LMXB). Based on the X-ray luminosity and the Bondi accretion rate, an IMBH accreting from the cluster gas seems unlikely and we suggest that the X-rays are due to accretion from a companion. Alternatively, the probability that a $1.5 M_\\odot$ cluster LMXB lies within the 95 per cent X-ray error circle is about 0.7. Therefore we cannot rule out a single LMXB as the origin of the X-ray emission. While we cannot distinguish between different models with current observations, future high-resolution and high-sensitivity radio imaging observations will reveal whether there is an IMBH at the centre of G1. ",
        "introduction": "Intermediate-mass black holes (IMBHs) have been a subject of debate for a long time. If they exist, IMBHs represent the long sought after link between stellar-mass and super-massive black holes. Until recently, we only have had indirect evidence for IMBHs via X-ray observations. For instance, ultraluminous ($L_X>10^{39}$\\lum) or hyperluminous ($L_X>10^{41}$\\lum) X-ray sources have been the best candidates (e.g. Wolter et al. 2006;  Farrell et al. 2009) based on their X-ray luminosities which, assuming isotropic emission, are far in excess of the Eddington limit for stellar mass black holes. Furthermore, X-ray spectroscopy and X-ray timing also provide some support for IMBHs (e.g. Miller et al. 2004; Kong \\& Di Stefano 2005; Strohmayer \\& Mushotzky 2009). However, we still require dynamical evidence in order to confirm solidly that IMBHs exist.  The only dynamical mass measurement claimed for an IMBH candidate is the globular cluster G1 in M31. G1 is the most luminous star cluster in the Local Group, and also one of the most massive at $(4\u2013-7)\\times 10^6\\,M_\\odot$ (Ma et al. 2009; Barmby et al. 2007). Based on Keck and {\\it Hubble Space Telescope (HST)} observations, it has been claimed that G1 hosts a $\\sim 2\\times10^4 \\, M_\\odot$ object at the core (Gebhardt et al. 2002,2005). However, the suggestion of an IMBH is controversial and has been challenged by Baumgardt et al. (2003). Recently, X-ray emission near the core of G1 has been discovered based on \\xmm\\ observations (Trudolyubov \\& Priedhorsky 2004; Pooley \\& Rappaport 2006; Kong 2007) and it is suggested that the X-rays come from Bondi accretion from cluster gas onto a central IMBH. However, it is also possible that the X-ray emission is due to an ordinary low-mass X-ray binary (LMXB), or a collection of faint LMXBs near the core. By refining the relative astrometry between \\xmm\\ and \\hst\\ data, Kong (2007) concluded that we cannot distinguish between these two scenarios. As we will show in this Letter, if there exists an IMBH in G1, accretion likely comes from a companion star (see, e.g. Patruno et a. 2006 for a scenario with a giant).  Alternatively, radio observations may be able to provide additional information about the nature of the X-ray source in G1. Ulvestad et al. (2007) employed the Very Large Array (VLA) to obtain a low-resolution ($\\sim 3''$ beam size at 8.4 GHz) image of G1 and a radio source was detected near the \\xmm\\ source, about an arcsecond from the cluster core. The radio/X-ray flux ratio of G1 ($\\sim 5\\times10^{-5}$) is at least a few hundred times higher than that expected for a LMXB near the cluster centre ($\\sim 5\\times10^{-8}$; see Fender \\& Kuulkers 2001), but is consistent with the expected value for accretion onto a $2\\times10^4 M_\\odot$ IMBH (Merloni et al. 2003). Furthermore, the radio/X-ray flux ratio is much lower than that of supernova remnants ($\\sim 10^{-2}$), but it is consistent with a pulsar wind nebula (Ulvestad et al. 2007).  However, it is still a question whether the X-ray and radio emission come from the same source because the resolution of both observations do not have sufficient accuracy to determine the precise position. While the VLA is still in its short baseline configuration and MERLIN is being upgraded, the first step is to use \\chandra\\ to obtain an accurate position for the X-ray source.  In this Letter, we localized the X-ray emission of G1 by performing precise relative astrometry using \\chandra, \\hst, and Canada-France-Hawaii Telescope (CFHT). In Section 2, we describe our X-ray and optical observations and data analysis. We present the localization of the X-ray source in Section 3. We finally discuss the nature of the X-ray source in Section 4. ",
        "conclusions": "By utilizing \\chandra, \\hst/WFPC2, and CFHT/MegaCam data, we determined the precise position of the X-ray emission of G1. The X-ray source, previously seen by \\xmm\\ (see Pooley \\& Rappaport 2006; Kong 2007), is very close ($\\sim 0\\farcs11$) to the cluster centre. Based on the calculation by Pooley and Rappaport (2006), if the X-ray emission is from Bondi accretion of ionized cluster gas by a central IMBH, the X-rays should come from the central 50 milli-arcsecond of the cluster. However, given the uncertainty of the error circle, we cannot rule out that the X-ray source is slightly offset from the cluster core (see Figure 1). If the X-ray emission is not from a central IMBH, it could come from a luminous LMXB located within or very close to the core radius ($0\\farcs21$ in Ma et al. 2007; $0\\farcs19$ in Barmby et al. 2007) of G1. As Kong (2007) points out, nearly half of the LMXBs in Galactic globular clusters are found within the core radius; it would not be surprising to find a luminous LMXB within the core of G1.   We estimated the probability of a $1.5 M_{\\odot}$ object being found in the X-ray error circle. The mass of $1.5 M_{\\odot}$ is the average quiescent LMXB mass, estimated from the radial distribution of 20 quiescent LMXBs in seven globular clusters (Heinke et al. 2003). We adopted the single-mass King model fit from Ma et al. (2007), which has $r_c = 0.21''$ and also considered the core radius value of $r_c = 0.19''$ from Barmby et al. (2007).  We assumed that the King model profile describes the distribution of turnoff-mass objects; we took the turnoff mass to be $0.9 M_\\odot$.  This gives a mass ratio of $q=1.67$ between the LMXBs and the turnoff-mass stars. It is straightforward to integrate the radial density profile for the LMXBs analytically.  As a check, half of the probability is within the core radius for $q=1.67$.  We then integrated the probability over angle numerically, to compute the probability within the off-centre X-ray error circle. We found that a $1.5 M_\\odot$ LMXB would have a probability of 0.756 ($r_c = 0.19''$) and 0.718 ($r_c = 0.21''$) of being found within our 95 per cent confidence \\chandra\\ error circle. Therefore the possibility that the X-ray source represents a single LMXB is significant.  If the X-ray source is from an accreting IMBH, the observed luminosity ($2\\times10^{36}$\\lum) and the spectrum would be consistent with a black hole in the hard state (Remillard \\& McClintock 2006).  There are also two possibilities if G1 has an accreting IMBH. The X-ray emission could be due to accretion either from cluster gas (Pooley \\& Rappaport 2006), or from a companion star.  We list in Table 1 a few examples of nearby supermassive black holes and dynamically confirmed stellar-mass black holes for comparison. All three quiescent supermassive black holes in the local universe have very low Eddington ratios of $10^{-11} - 10^{-9}$. On the other hand, quiescent stellar-mass black holes tend to have a higher Eddington ratio. V404 Cyg is the most X-ray luminous stellar-mass black hole while A0620--00 and GS2000+25 are the faintest ones (with detection). The Eddington ratio of G1 (assuming a mass of $2\\times10^4 M_\\odot$) is at least two orders of magnitude higher than that of quiescent supermassive black holes, but is in the range of stellar-mass black holes. We also compare the  X-ray luminosity in units of Bondi luminosity. This is a better indicator of accretion luminosity than the Eddington ratio, as it relates the X-ray luminosity to the available mass transfer rate. For G1, it is about 0.01 (see Pooley \\& Rappaport 2006), which is substantially higher than quiescent supermassive black holes (see Table 1). It is an indication that the accretion efficiency of G1 must be high if cluster gas is the source of the accretion.  In the framework of Bondi accretion, we further derive the Eddington-scaled Bondi accretion rate (see Table 1) and compare with other low luminosity supermassive black holes as shown in Figure 14 of Soria et al. (2006). According to the standard advection-dominated accretion flow (ADAF) model, the X-ray luminosity of G1 should be two orders of magnitude lower. This implies that the putative IMBH of G1 is extremely efficient in accreting cluster gas. It is worth noting that apart from 47 Tuc (Freire et al. 2001), we do not have evidence of cluster gas in Galactic globular clusters (van Loon et al. 2006).  Furthermore, if ionised cluster gas (as described in Pooley and Rappaport 2006) is the only source of inflow, such a high accretion rate would be difficult to reconcile with the ADAF prediction. It is therefore indicative that if the X-ray emission of G1 is from an IMBH, it is more likely due to accretion from a companion star. However, based on the current \\chandra\\ data, we cannot distinguish between the two possible mechanisms (IMBH or ordinary LMXB) for generating the X-ray emission of G1.  We can rule out that the X-ray emission is from combination of low luminosity sources. The total of all low luminosity sources in 47 Tuc is only $3.7\\times10^{33}$\\lum, implying that the encounter rate of G1 must be $> 500$ times of 47 Tuc. However, Pooley \\& Rappaport (2006) estimated that the encounter rate of G1 is only 17 times of 47 Tuc. Hence the X-ray emission must come from one or a few bright sources. Supernova remnants may also be responsible for the X-ray emission but the probability of a core collapse supernova in a globular cluster is extremely small, as is the probability of a Type Ia supernova (e.g. Pfahl et al. 2009). Indeed, if the X-ray source is associated with the radio source detected with the VLA (Ulvestad et al. 2007), both the LMXB and supernova remnant scenarios are unlikely based on the radio/X-ray flux ratio (see below). The probability of positional coincidence between G1 and an unrelated X-ray source is very small. We estimated the background AGN/foreground star rate within the half-mass radius of G1 ($\\sim 1''$; Barmby et al. 2007) using the \\chandra\\ Deep Field (Brandt et al. 2001). At an 0.5--2 keV flux of $1.2\\times10^{-14}$ erg s$^{-1}$ cm$^{-2}$, we expect to find on average $4\\times10^{-5}$ background or foreground objects within the half mass radius.  Using the VLA, Ulvestad et al. (2007) detected a radio source coincident with G1 at 8.4 GHz with a flux density of $28\\pm6 \\mu$Jy. However, the radio position is offset from the X-ray and optical centre by about $1\\farcs3$ (see Figure 1) using our refined \\chandra\\ and \\hst\\ observations. Because of the short baseline (3.5 km) of the VLA observations, the rms error of the VLA position is about $0\\farcs6$ in each dimension. Therefore, the cluster core is still within the 95 per cent error radius (see Figure 1). In order to confirm if the radio source is associated with the X-ray source, we will require high-resolution and high-sensitivity radio observations using the Expanded VLA (EVLA) in its long baseline configuration, e-MERLIN, or VLBI.  If we can confirm that the radio source is related to the X-ray source, this will provide a strong support that there is an IMBH near the centre of G1 because the radio/X-ray flux ratio of G1 ($\\sim 5\\times10^{-5}$) is several hundreds times higher than that of a LMXB, but is consistent with a $2\\times10^4 M_\\odot$ IMBH using the relationship in the ``fundamental plane'' linking radio/X-ray flux ratios and black hole masses (Merloni et al. 2003; see also Table 1). The radio/X-ray flux ratio of G1 is also substantially lower than that of supernova remnants ($\\sim 10^{-2}$) and low-luminosity active galactic nuclei (see Ulvestad et al. 2007 and Table 1). The radio/X-ray flux ratio is consistent with a pulsar wind nebula, though we have no evidence for pulsars with such high rotational energy losses having been born in old populations such as G1.  A radio spectrum (with EVLA or e-MERLIN) or milliarcsecond-resolution VLBI observations, as suggested by Ulvestad et al. (2007), could rule this possibility out, as could detection of strong X-ray variability. It is also possible that the radio source is not coincident with the X-ray source. In this case, the radio source may be a background source, a radio-loud X-ray binary in G1, or a jet from the central IMBH."
    },
    "0910/0910.1200_arXiv.txt": {
        "abstract": "In the 5-year WMAP data analysis, a new parametrization form for dark energy equation-of-state was used, and it has been shown that the equation-of-state, $w(z)$, crosses the cosmological-constant boundary $w=-1$. Based on this observation, in this paper, we investigate the reconstruction of quintom dark energy model. As a single-real-scalar-field model of dark energy, the generalized ghost condensate model provides us with a successful mechanism for realizing the quintom-like behavior. Therefore, we reconstruct this scalar-field quintom dark energy model from the WMAP 5-year observational results. As a comparison, we also discuss the quintom reconstruction based on other specific dark energy ansatzs, such as the CPL parametrization and the holographic dark energy scenarios. ",
        "introduction": "\\label{sec:intr}  Observations of high-redshift supernovae indicate that the universe is accelerating at the present stage~\\cite{SN} and this accelerating expansion also has been confirmed by many other cosmological experiments, such as observations of large scale structure (LSS) \\cite{LSS}, and measurements of the cosmic microwave background (CMB) anisotropy \\cite{CMB}. We refer to the cause for this cosmic acceleration as ``dark energy,'' which is a mysterious exotic matter with large enough negative pressure and whose energy density has been a dominative power of the universe. The combined analysis of cosmological observations suggests that the universe is consists of about $70\\%$ dark energy, $30\\%$ dust matter (cold dark matter plus baryons), and negligible radiation. The astrophysical feature of dark energy is that it remains unclustered at all scales where gravitational clustering of baryons and nonbaryonic cold dark matter can be seen. Its gravity effect is shown as a repulsive force so as to make the expansion of the universe accelerate when its energy density becomes dominative power of the universe.  Although the nature and origin of dark energy are unknown, we still can propose some candidates to describe the properties of dark energy. The most obvious theoretical candidate of dark energy is the cosmological constant $\\Lambda$ (vacuum energy) \\cite{Einstein:1917} with an equation of state $w=-1$. The cosmological constant is rather popular in researches of cosmology and astrophysics due to its theoretical simpleness and its great success in fitting with observational data. However, as is well known, the two fundamental problems, namely the ``fine-tuning'' problem and the ``cosmic coincidence'' problem \\cite{coincidence}, still puzzle us. Theorists have made many efforts to try to resolve the cosmological constant problem, but all of these efforts turn out to be unsuccessful \\cite{cc}.  Also, there are other alternatives to the cosmological constant. An alternative proposal to explaining dark energy is the dynamical dark energy scenario. The dynamical dark energy proposal is often realized by some scalar field mechanism which suggests that the energy form with negative pressure is provided by a scalar field slowly rolling down its potential. So far, a lot of scalar-field dark energy models have been studied. The models such as quintessence \\cite{quintessence}, $K$-essence \\cite{kessence}, phantom \\cite{phantom}, tachyon \\cite{tachyon} and ghost condensate \\cite{ghost1,ghost2} are all famous examples of scalar-field dark energy models. In these models, the quintessence with a canonical kinetic term evolves its equation of state in the region of $w\\geqslant -1$ whereas the model of phantom with negative kinetic term can always lead to $w\\leqslant -1$; the $K$-essence can realize both $w>-1$ and $w<-1$, but it has been shown that it is very difficult for $K$-essence to achieve $w$ of crossing $-1$ \\cite{Vikman:2004dc}.  However, the analysis of the current observational data shows that the equation of state of dark energy $w$ is likely to cross the cosmological-constant boundary (or phantom divide) $-1$, i.e. $w$ is larger than $-1$ in the recent past and less than $-1$ today. The dynamical evolving behavior of dark energy with $w$ getting across $-1$ has brought forward great challenge to the model-building of scalar-field in the cosmology. Just as mentioned above, the scalar-field models, such as quintessence, $K$-essence, phantom, cannot realize the transition of $w$ from $w>-1$ to $w<-1$ or vice versa. Hence, the quintom model was proposed for describing the dynamical evolving behavior of $w$ crossing $-1$ \\cite{quintom} with double fields of quintessence and phantom. The cosmological evolution of such model has been investigated in detail \\cite{twofield,quintom2}. For the single real scalar field models, the transition of crossing $-1$ for $w$ can occur for the Lagrangian density $p(\\phi, X)$, where $X$ is a kinematic term of a scalar-field $\\phi$, in which $\\partial{p}/\\partial{X}$ changes sign from positive to negative, thus we require nonlinear terms in $X$ to realize the $w=-1$ crossing \\cite{ghost2,Vikman:2004dc,Anisimov:2005ne}. When adding a high derivative term to the kinetic term $X$ in the single scalar field model, the energy-momentum tensor is proven to be equivalent to that of a two-field quintom model \\cite{Li:2005fm}. It is remarkable that the generalized ghost condensate model of a single real scalar field is a successful realization of the quintom-like dark energy \\cite{Tsujikawa:2005ju,Zhang:2006qu}. What's more, a generalized ghost condensate model was investigated in Refs.~\\cite{Tsujikawa:2005ju,Zhang:2006qu} by means of the cosmological reconstruction program. For another interesting single-field quintom model see Ref.~\\cite{Huang:2005gu}, where the $w=-1$ crossing is implemented with the help of a fixed background vector field. Besides, there are also many other interesting models, such as holographic dark energy model \\cite{holo} and braneworld model \\cite{Cai:2005ie}, being able to realize the quintom-like behavior.  In any case, these dark energy models including the dynamical dark energy models have to face the test of cosmological observations. A typical approach for this is to predict the cosmological evolution behavior of the models, by putting in the Lagrangian (in particular the potential) by hand or theoretically, and to make a consistency check of models by comparing it with observations. An alternative approach is to reconstruct corresponding theoretical Lagrangian, by using the observational data. The reconstruction of scalar-field dark energy models has been widely studied. For a minimally coupled scalar field with a potential $V(\\phi)$, the reconstruction is simple and straightforward \\cite{simplescalar}. Saini et al. \\cite{Saini:1999ba} reconstructed the potential and the equation of state of the quintessence field by parameterizing the Hubble parameter $H(z)$ based on a versatile analytical form of the luminosity distance $d_L(z)$. This method can be generalized to a variety of models, such as scalar-tensor theories \\cite{scalartensor}, $f(R)$ gravity \\cite{frgrav}, $K$-essence model \\cite{Li:2006bx,Gao:2007ep}, and also tachyon model \\cite{holotach}, etc.. Tsujikawa has investigated the reconstruction of general scalar-field dark energy models in detail \\cite{Tsujikawa:2005ju}.  In this paper, we will investigate the quintom reconstruction from the Wilkinson Microwave Anisotropy Probe (WMAP) 5-year observations. We will focus on the generalized ghost condensate model and will reconstruct this quintom scalar-field model using various dark energy ansatzs including the parametric forms of dark energy and holographic dark energy scenarios. In particular, we will put emphasis on a new parametrization form proposed by WMAP team in Ref.~\\cite{Komatsu:2008hk}.  The paper is organized as follows: In section \\ref{sec:para} we address the dark energy parametrization proposed in Ref.~\\cite{Komatsu:2008hk} and describe the corresponding analysis results of the WMAP5 observations. In section \\ref{sec:ghost} we perform a cosmological reconstruction for the generalized ghost condensate model from various dark energy ansatzs and the fitting results of the up-to-date observational data. Finally we give the concluding remarks in section \\ref{sec:concl}. ",
        "conclusions": "\\label{sec:concl}  The recent fits to current observational data, such as SN, CMB and LSS, find that even though the behavior of dark energy is consistent to great extent with a cosmological constant, an evolving dark energy with the equation of state $w$ larger than $-1$ in the recent past but less than $-1$ today is also with some possibility. Although the scalar-field models of dark energy, such as quintessence and phantom, can provide us with dynamical mechanism for dark energy, the behavior of cosmological-constant crossing brings forward a great challenge to the model-building for dynamical dark energy, because neither quintessence nor phantom can fulfill this behavior. A two-field quintom model, therefore, was suggested to realize this behavior by means of the incorporation of the features of quintessence and phantom. Besides, the generalized ghost condensate model provides us with a successful single-real-scalar-field model for realizing the quintom-like behavior. For probing the dynamical nature of dark energy, one should parameterize dark energy first and then constrain the parameters using the observational data.  In this paper, we have investigated the dynamical behavior of the general ghost condensate scalar-field model by taking a new form of parametrization of the equation of state proposed in Ref.~\\cite{Komatsu:2008hk} and the best-fit values from the observational data. The results of reconstruction show the dynamical behavior of the generalized ghost condensate from this parametrization. In particular, this reconstruction investigation explores the early-time evolutionary behavior of the scalar field model. As a comparison, we also discussed other specific cases including the CPL parametrization and the holographic dark energy scenarios.   The increase of the quantity and quality of observational data in the future will undoubtedly provide a true {\\it model-independent} manner for exploring the properties of dark energy. We hope that the future high-precision observations (e.g. SNAP) may be capable of providing us with deep insight into the nature of dark energy driving the acceleration of the universe."
    },
    "0910/0910.4339_arXiv.txt": {
        "abstract": "% With the ever-growing number of exoplanets detected, the issue of characterization is becoming more and more relevant. Direct imaging is certainly the most efficient but the most challenging tool to probe the atmosphere of exoplanets and hence in turns determine the physical properties and refine models of exoplanets. A number of instruments optimized for exoplanets imaging are now operating or planned for the short and long term both on the ground and in space. This paper reviews these instruments and their characteristics/capabilities. Conclusions are drawn on the spectral characterization point of view. ",
        "introduction": "The study of extrasolar planets has became in a decade an exciting field in modern astronomy. With more than 300 planets detected so far, indirect methods have been the most prolific at finding sub-stellar objects in the solar neighborhood. However, a few angularly resolved images of Extrasolar Giant Planets (EGPs) have been finally obtained with 8-10m class telescopes on the ground as well as in space with the HST \\citep[for instance,][]{Chauvin05, Neuhauser05, Lafreniere08}. The detection rate is progressing quickly. Some discoveries were announced a week or so before the conference \\citep{Kalas08, Marois08, Lagrange08}. These outstanding observations were made possible by the favorable configurations of these planetary systems: low star/planet mass ratio, young age (a few tens to hundreds Myr), large physical distances (typically hundreds of AU, except for $\\beta$ Pic) which basically make these planets bright enough or locally above the stellar halo. A list of objects detected by direct imaging is summarized in Tab. 1.  This paper presents an overview of direct imaging projects on single apertures. As guideline for non-specialists, the discussion remains general and more details can be found in reference papers. The first part briefly reminds the problematic of high contrast imaging and the instrumental solutions that are being developed. In section \\ref{sec:planned} and \\ref{sec:future} the projects are separated in two categories: those that are planned and those that are proposed for the future and mostly focused on telluric planets. The intent is to clearly points out what is going to be achieved in terms of parameter space in order to better define what will be needed for the next projects. Finally, all these projects are sorted along a timeline and some conclusions are derived.  \\begin{table}[!ht] \\caption{List of planetary mass objects detected with direct imaging, as of Feb. 2009. Error bars are not indicated but are often larger than 1\\,M$_J$ therefore placing some of these objects in the Brown Dwarf regime.} \\smallskip \\begin{center} {\\small \\begin{tabular}{lccl} \\tableline \\noalign{\\smallskip} object \t&\testimated mass &  estimated separation & reference \\\\ & \t\t[M$_J$]\t &\t[AU] \t\t\t\t& \\\\ \\noalign{\\smallskip} \\tableline \\noalign{\\smallskip} 2M1207 b\t\t& \t5 \t\t&\t46   & \\citet{Chauvin05} \\\\ GQ Lup b\t\t& \t17 \t\t& \t100 & \\citet{Neuhauser05}\\\\ AB Pic b\t\t& \t14 \t\t&\t248 & \\citet{Chauvin05b}\\\\ CHRX73 b\t& \t12 \t\t&\t210 & \\citet{Luhman06}\\\\ HN Peg b\t\t& \t16 \t\t&\t795 & \\citet{Luhman07} \\\\ DH Tau b\t\t& \t12 \t\t&\t330 & \\citet{Itoh05}\\\\\\ RSX 1609 b\t& \t8\t\t&\t330 & \\citet{Lafreniere08} \\\\ Fomalhaut b\t& \t3  \t\t&\t120 & \\citet{Kalas08}\\\\ HR8799 b\t\t& \t5 \t\t& \t46   & \\citet{Marois08}\\\\ HR8799 c\t\t& \t12 \t\t& \t330 & \\citet{Marois08}\\\\ HR8799 d\t\t& \t8\t\t&  \t330 & \\citet{Marois08}\\\\ $\\beta$ Pic b\t&\t8\t\t&\t8      & \\citet{Lagrange08} \\\\ \\noalign{\\smallskip} \\tableline \\end{tabular} } \\end{center} \\end{table} ",
        "conclusions": "From the above description, it is instructive to put on a timeline the different planned projects identifying spectral range and typical detectable planets. It is obviously recognized that all bandwidths are of interest for the characterization of exoplanets. For the direct imaging, this applies to visible, near IR and mid IR. What comes up from Fig. 1 is that a slot for a space mission geared toward mature giant planets either at visible or mid IR is not covered on the ground at least until 2025-2030. Another way of putting this information in a mass vs. wavelength diagram is shown in Fig. 2. If we apply a concept of minimal effort, this slot could be fill in by small telescopes optimized for a deep characterization of giant planets at visible an ideally down to  massive telluric planets that are named Super Earths. If all telluric planets share similar characteristics (atmospheric composition, geophysical attributes, biosignatures, seasonnal activities, ...), it would be worth to devote a space mission on that topic because it would be both complementary to ground-based projects and also preparatory to those that are focused to the search of life on Earth analogs \\citep{cockell08}."
    },
    "0910/0910.4992.txt": {
        "abstract": "Although the nature of the progenitor of Type Ia supernovae (SNe Ia) is still unclear, the single-degenerate (SD) channel for the progenitor is currently accepted, in which a carbon-oxygen white dwarf (CO WD) accretes hydrogen-rich material from its companion, increases its mass to the Chandrasekhar mass limit, and then explodes as a SN Ia. The companion may be a main sequence or a slightly evolved star (WD + MS), or a red-giant star (WD + RG). Incorporating the effect of mass-stripping and accretion-disk instability on the evolution of WD binary, we carried out binary stellar evolution calculations for more than 1600 close WD binaries. As a result, the initial parameter spaces for SNe Ia are presented in an orbital period-secondary mass ($\\log P_{\\rm i}, M_{\\rm 2}^{\\rm i}$) plane. We confirmed that in a WD + MS system, the initial companion leading to SNe Ia may have mass from 1 $M_{\\odot}$ to 5 $M_{\\odot}$. The initial WD mass for SNe Ia from WD + MS channel is as low as $0.565$ $M_{\\odot}$, while the lowest WD mass from the WD + RG channel is 1.0 $M_{\\odot}$. Adopting the results above, we studied the birth rate of SNe Ia via a binary population synthesis approach. We found that the Galactic SNe Ia birth rate from SD model is $2.55 - 2.9\\times10^{\\rm -3}$ ${\\rm yr}^{\\rm -1}$ (including WD + He star channel), which is slightly smaller than that from observation. If a single starburst is assumed, the distribution of the delay time of SNe Ia from the SD model may be a weak bimodality, where WD + He channel contributes to SNe Ia with delay time shorter than $10^{\\rm 8}$ yr and WD + RG channel to those with age longer than 6 Gyr. ",
        "introduction": "\\label{sect:1} \\subsection{Progenitor Models}\\label{sect:1.1} Type Ia supernovae (SNe Ia) play an important role in astrophysics, especially in cosmology. They appear to be good cosmological distance indicators and are successfully applied in determining cosmological parameters (e.g. $\\Omega$ and $\\Lambda$), which leads to the discovery of the accelerating expansion of the universe (\\citealt{RIE98}; \\citealt{PER99}). At present, SNe Ia are proposed to be cosmological probes for testing the evolution of the Dark Energy equation of state with time (\\citealt{HOWEL09}). However, the progenitor systems of SNe Ia have not yet been confidently identified (\\citealt{HN00}; \\citealt{LEI00}), although they show their importance in many astrophysical field, e.g. determining cosmological parameters, studying galaxy evolution, understanding explosion mechanism of SNe Ia and constraining the theory of binary stellar evolution (see the review by \\citealt{LIVIO99}).  There is a consensus that SNe Ia result from the explosion of a carbon-oxygen white dwarf in a binary system (\\citealt{HF60}). The CO WD accretes material from its companion, increases mass to its maximum stable mass, and then explodes as a thermonuclear runaway. Almost half of the WD mass is converted into radioactive nickel-56 in the explosion (\\citealt{BRA04}), and the amount of nickel-56 determines the maximum luminosity of SNe Ia (\\citealt{ARN82}). According to the nature of the companions of the mass accreting white dwarfs, two basic scenarios for the progenitor of SN Ia have been discussed over the last three decades. One is the single degenerate (SD) model (\\citealt{WI73}; \\citealt{NTY84}), in which a CO WD increases its mass by accreting hydrogen- or helium-rich matter from its companion, and explodes when its mass approaches the Chandrasekhar mass limit. The companion may be a main-sequence or a slightly evolved star (WD+MS) or a red-giant star (WD+RG) or a helium  star (WD + He star) (\\citealt{YUN95}; \\citealt{LI97}; \\citealt{HAC99a,HAC99b}; \\citealt{NOM99, NOM03}; \\citealt{LAN00}; \\citealt{HAN04, HAN06}; \\citealt{CHENWC07, CHENWC09}; \\citealt{HAN08}; \\citealt{MENG09}; \\citealt{MENGYANG09}; \\citealt{LGL09}; \\citealt{WANGB09a, WANGB09b}). An alternative is the double degenerate (DD) model (\\citealt{IT84}; \\citealt{WEB84}), in which a system consisting of two CO WDs loses orbital angular momentum by gravitational wave radiation and merges finally. The merger may explode if the total mass of the system exceeds the Chandrasekhar mass limit (see the reviews by \\citealt{HN00} and \\citealt{LEI00}). Although the channel is theoretically less favored, e.g. double WD mergers may lead to accretion-induced collapses rather than to SNe Ia (\\citealt{HN00}), it is premature to exclude the channel at present since there exists evidence that the channel may contribute to a few SNe Ia (\\citealt{HOW06}; \\citealt{BRA06}; \\citealt{QUI07}; \\citealt{HICKEN07}; \\citealt{YUAN07}). At present, the single-degenerate Chandrasekhar model is widely accepted, since it is supported by many observations (\\citealt{PAR07})  \\subsection{Observations}\\label{sect:1.2} The SD model is supported by many observations. For example, variable circumstellar absorption lines were observed in the spectra of SN Ia 2006X (\\citealt{PAT07}), which indicates the SD nature of its precursor. \\citet{PAT07} suggested that the progenitor of SN 2006X is a WD + RG system based on the expansion velocity of the circumstella material, while \\citet{HKN08} argued a WD + MS nature for this SN Ia. Recently, two twins of SN 2006X were also reported (\\citealt{BLONDIN09}; \\citealt{SIMON09}), and the birth rate and age of the 2006X-like supernovae may be well explained by the WD + MS model (\\citealt{MYG09a}). \\citet{VOSS08} suggested that SN 2007on is also possibly from a WD + MS channel. Moreover, several SD systems are possible progenitors of SN Ia. For example, supersoft X-ray sources (SSSs) were suggested as good candidates for the progenitors of SNe Ia (\\citealt{LIVIO96}; \\citealt{KV97}; \\citealt{HK03a, HK03b}). Some of the SSSs are WD + MS systems and some are WD + RG systems (\\citealt{DIK03}). A direct way to confirm the progenitor model is to search for the companion stars of SNe Ia in their remnants. The discovery of the potential companion of Tycho's supernova may have verified the reliability of the WD + MS model (\\citealt{RUI04}; \\citealt{BRA04}; \\citealt{IHA07}; \\citealt{HERNANDEZ09}).  Measuring the delay time of SNe Ia (DD, between the episode of star formation producing progenitor systems and the occurrence of SNe Ia) is a very important way to constrain the progenitor system. Although the DD of most of SNe Ia is between 0.3-2 Gyr (\\citealt{SCHAWINSKI09}), there are still SNe Ia with long delay times (older then 10 Gyr inferred from SNe Ia in elliptical galaxies in the local universe, \\citealt{MAN05}) or extremely short delay times (shorter than 0.1-0.3 Gyr, \\citealt{MAN06}; \\citealt{SCHAWINSKI09}; \\citealt{RASKIN09}). Some observational results, i.e. the strong enhancement of the SN Ia birthrate in radio-loud early-type galaxies, the strong dependence of the SN Ia birthrate on the colors of the host galaxies, and the evolution of the SN Ia birthrate with redshift (\\citealt{DEL05}; \\citealt{MAN05, MAN06}), may indicate a bimodal distribution of delay time (DDT), in which a part of the SNe Ia explode soon after starburst with a delay time less than 0.1-0.3 Gyr (`prompt' SNe Ia), while the rests have a much wider distribution with a delay time of about 3 Gyr (`tardy' SNe Ia \\citealt{MAN06}). 10\\% (weak bimodality) to 50\\% (strong bimodality) of all SNe Ia belong to the prompt SNe Ia (\\citealt{MANNUCCI08}). Whether this is due to the presence of two different channels of explosion (such as single- and double-degenerate) or to a bimodal distribution of some parameter of the exploding systems (such as the mass ratio between the two stars of the binary system) is still unclear.  \\subsection{Situations}\\label{sect:1.3} Many works have concentrated on the SD model. To explain prompt SNe Ia, various models were designed. \\citet{HKN08} introduced a mass-stripping effect on a main-sequence (MS) or slightly evolved companion star by winds from a mass-accreting white dwarf (\\citealt{HAC96}). \\citet{WANGB09a, WANGB09b} investigated WD + He star channel which can explain SNe Ia with short delay times very well. \\citet{LGL09} even designed a symbiotic channel by assuming an aspherical stellar wind with an equatorial disk. For the tardy SNe Ia, most of works focus on the WD + MS and WD + RG channel. Some authors (\\citealt{HAC99a, HAC99b, HKN08}; \\citealt{NOM99, NOM03}) have studied the SD channel by a simple analytical method to treat binary interactions. Such analytic prescriptions may not describe some mass-transfer phases, especially those occurring on a thermal time-scale (\\citealt{LAN00}). \\citet{LI97} studied the SD from detailed binary evolution calculation, while they considered two WD masses (1.0 and 1.2 $M_{\\odot}$). \\citet{LAN00} investigated the channel for metallicities $Z=0.001$ and 0.02, but they only studied the case A evolution (mass transfer during core hydrogen burning phase). \\citet{HAN04} carried out a detailed study of this channel including case A and early case B (mass transfer occurs at Hertzsprung gap, HG) for $Z=0.02$. Following the study of \\citet{HAN04}, \\citet{MENG09} studied the WD + MS channel comprehensively and systematically at various Z. Based on the works above, WD + MS channel is a very important channel to produce SNe Ia, but the channel may only account for about 1/3 SNe Ia observed and the delay time for the channel is between 0.2-2 Gyr (\\citealt{HAN04}; \\citealt{MENG09}). Recently, \\citet{XL09} proposed a WD + MS model with a thermal instable accretion disk, which can also produce SNe Ia with long delay time. Although WD + RG channel has been widely studied, it should be investigated carefully by detailed binary evolution calculations with the latest input physics since either the studies are based on a simple analytical method or focus on some special WD masses. In this paper, we want to construct a complete model, which includes WD + MS and WD + RG channels.  In Section \\ref{sect:2}, we describe binary evolution model, and present the calculation results in Section \\ref{sect:3}. We describe our binary population synthesis (BPS) method in Section \\ref{sect:4} and show the BPS results in Section \\ref{sect:5}. We give discussions in Section \\ref{sect:6} and summarize our main conclusions in Section \\ref{sect:7}.  \\begin{figure*} \\centerline{\\includegraphics[angle=270,scale=.7]{lummdot1.ps}} %\\centerline{\\psfig{figure=lummdot1.ps,width=14.cm,height=160mm,bbllx=60pt,bblly=100pt,bburx=585pt,bbury=676pt,angle=270}} \\caption{Three representative examples of binary evolution calculations. The solid, dashed, dash-dotted and dotted curves show the mass-transfer rate, $\\dot{M}_{\\rm tr}$, the mass-growth rate of the CO WD, $\\dot{M}_{\\rm WD}$, the mass of the CO WD, $M_{\\rm WD}$, and stripping-off mass-lose rate from secondary, $\\dot{M}_{\\rm strip}$, respectively, in panels (b), (d) and (f). Dotted vertical lines in all panels and asterisks in panels (a), (c) and (e) represent the position where the WD is expected to explode as a SN Ia.}\\label{lummdot1} \\end{figure*} ",
        "conclusions": "\\label{sect:7} Incorporating mass-stripping effect in \\citet{HKN08} and the effect of the instability of accretion disk on the evolution of WD binaries into Eggleton's stellar evolution code, and including also the prescription of \\citet{HAC99a} for mass-accretion of CO WD, We performed binary evolution calculations for more than 1600 close WD binaries. Adopting the results obtained here, we carried out a series BPS sutdy and calculated the birth rate of SNe Ia. We summarize the basic results as following.  1. The detailed binary evolution calculation results further confirmed the suggestion in \\citet{HKN08} that the mass-stripping effect may attenuate between WD and the companion, and avoid the formation of a CE, and then increase the donor mass leading to SNe Ia. For a reasonable strength of the mass-stripping effect in physics, i.e. $c_{\\rm 1}=1.5$, a companion as massive as 5 $M_{\\odot}$ can produce an SNe Ia. However, the detailed BPS study show that the upper limit of the companion could be $\\sim4.5 M_{\\odot}$.  2. The detailed binary evolution calculation results confirm that the disk instability could substantially increase the mas-accumulation efficiency for accreting WD, and extend the mass of donor star producing SNe Ia to $\\sim 1 M_{\\odot}$ (see also \\citealt{XL09} and \\citealt{WANGB09c}).  3. CO WDs leading to SNe Ia may have a mass as low as $0.565 M_{\\odot}$.  4. The Galactic birth rate from the SD channel in this paper is $2.25-2.6\\times10^{\\rm -3}$ ${\\rm yr}^{\\rm -1}$, based on $\\alpha_{\\rm CE}$. If the WD + He star channel is also included, the Galactic birth rate from the SD channel increases to $2.55-2.9\\times10^{\\rm -3}$ ${\\rm yr}^{\\rm -1}$, only slightly smaller than that derived from observation. Then, there should be other channels or mechanisms contributing to SNe Ia.  5. The DDT from the WD + MS, WD + RG and WD + He star channels may be a weak bimodality. The WD + MS produces most of SNe Ia with delay time between $0.15-2.5$ Gyr. The WD + RG channel may mainly contribute to SNe Ia with delay time longer than 6 Gye.  6. The mass-stripping effect may not contribute to very young population of SNe Ia. The WD + He star channel could contribute most of young SNe Ia, even though not all.  7. Our BPS results may uphold the C/O as the origin of the luminosity scatter of SNe Ia and provide a method to verify it.  8. We show the evolution of the initial parameters of SD systems for SNe Ia with time, which is helpful for searching for possible candidates of SNe Ia progenitor systems.   %% The equation environment wil produce a numbered display equation.    %% The \\notetoeditor{TEXT} command allows the author to communicate %% information to the copy editor.  This information will appear as a %% footnote on the printed copy for the manuscript style file.  Nothing will %% appear on the printed copy if the preprint or %% preprint2 style files are used.  %% The eqnarray environment produces multi-line display math. The end of %% each line is marked with a \\\\. Lines will be numbered unless the \\\\ %% is preceded by a \\nonumber command. %% Alignment points are marked by ampersands (&). There should be two %% ampersands (&) per line.    %% Putting eqnarrays or equations inside the mathletters environment groups %% the enclosed equations by letter. For instance, the eqnarray below, instead %% of being numbered, say, (4) and (5), would be numbered (4a) and (4b). %% LaTeX the paper and look at the output to see the results.    %% This section contains more display math examples, including unnumbered %% equations (displaymath environment). The last paragraph includes some %% examples of in-line math featuring a couple of the AASTeX symbol macros.    %% The displaymath environment will produce the same sort of equation as %% the equation environment, except that the equation will not be numbered %% by LaTeX.      %% If you wish to include an acknowledgments section in your paper, %% separate it off from the body of the text using the"
    },
    "0910/0910.2804_arXiv.txt": {
        "abstract": "{}{ We investigate the size--density relation in extragalactic \\hii\\ regions, with the aim of understanding the role of dust and different physical conditions in the ionized medium. } {First, we compiled several observational data sets for Galactic and extragalactic \\hii\\ regions and confirm that extragalactic \\hii\\ regions follow the same size ($D$)--density ($n$) relation as Galactic ones ($n\\propto D^{-1}$), rather than a relation with constant luminosity ($n\\propto D^{-1.5}$). Motivated by the inability of static models to explain this, we then modelled the evolution of the size--density relation of \\hii\\ regions by considering their star formation history, the effects of dust, and  pressure-driven expansion. The results are compared with our sample data whose size and density span roughly six orders of magnitude. } {The extragalactic samples cannot be understood as an evolutionary sequence with a single initial condition. Thus, the size--density relation does not result from an evolutionary sequence of \\hii\\ regions but rather reflects a sequence with different initial gas densities (``density hierarchy''). We also find that the size of many \\hii\\ regions is limited by dust absorption of ionizing photons, rather than consumption by ionizing neutral hydrogen. Dust extinction of ionizing photons is particularly severe over the entire lifetime of compact \\hii\\ regions with typical gas densities of $\\ga 10^3$ cm$^{-3}$. Hence, as long as the number of ionizing photons is used to trace massive star formation, much star-formation activity could be missed. Such compact dense environments, the ones most profoundly obscured by dust, have properties similar to ``maximum--intensity starbursts''. This implies that submillimeter and infrared wavelengths may be necessary to accurately assess star formation in these extreme conditions both locally and at high redshift. } {} ",
        "introduction": "\\label{sec:intro}  Massive stars, young star clusters, and their associated \\hii\\ regions are vital probes of recent star formation in nearby galaxies and the distant universe. While Galactic \\hii\\ regions such as the Orion Nebula and RCW\\,49 are well studied because of their proximity, most Galactic work is still plagued by distance uncertainties (e.g., Anderson \\& Bania \\cite{anderson09}), although the situation is improving rapidly (see, e.g., Hachisuka et al. \\cite{hachisuka06}, Foster \\& MacWilliams \\cite{foster06}, Russeil et al. \\cite{russeil07}). \\hii\\ regions in galaxies in the Local Group can be studied in almost as much detail as those in the Galaxy, and the distance estimation is much less uncertain. However, even in Local Group galaxies, Super Star Clusters (SSCs), the most extreme examples of massive star formation, are absent. SSCs, with $\\sim10^5$--$10^6$\\,\\msun\\ and Lyman continuum photon rates $\\Nion\\sim 10^{52}$--$10^{53}$\\,s$^{-1}$ enclosed in regions of $\\la$10\\,pc in radius, are generally not seen in quiescent environments such as the Milky Way, M\\,31, and M\\,33. Even R136 in 30 Doradus, the most massive Local Group star cluster (10$^{4.5}$\\,\\msun\\ and 10$^{51.4}$\\,s$^{-1}$: Massey \\& Hunter \\cite{massey98}), falls short of the typical properties of SSCs. Hence, to study the wide range of manifestations of massive star formation, it is necessary to examine star clusters and \\hii\\ regions in galaxies beyond the Local Group. Of necessity, the price to be paid is detail; the advantage to be gained is the wide variety of Star-Forming (SF) complexes that can be studied.  The total mass of a cluster of massive stars can be observationally quantified by estimating the number of ionizing photons. A simple Str\\\"{o}mgren-sphere argument (Eq.\\ \\ref{eq:stromgren}) would suggest that \\nion$\\propto D^3n_\\mathrm{e}^2$, where \\nion\\ is the number of ionizing photons emitted per unit time, $D$ is the diameter of the \\hii\\ region, and $n_\\mathrm{e}$ is the electron number density in the \\hii\\ region. Thus, in principle, by examining $n_\\mathrm{e}$ and $D$ of \\hii\\ regions, we can estimate the ionizing photon luminosity. However, observationally, the situation is not straightforward. The diameter ($D$) and the electron density ($n_\\mathrm{e}$) of Galactic \\hii\\ regions are known to be negatively correlated with a roughly unit slope: $n_\\mathrm{e}\\propto D^{-1}$ (Garay \\& Lizano \\cite{garay99}; Kim \\& Koo \\cite{kim01}; Mart\\'{\\i}n-Hern\\'{a}ndez et al.\\ \\cite{martin03}; Dopita et al.\\ \\cite{dopita06}). This would not be expected from a simple Str\\\"{o}mgren-sphere argument with \\nion\\ constant: $n_\\mathrm{e}\\propto D^{-1.5}$. Various possible explanations for the observed shallower slope are presented in Mart\\'{\\i}n-Hern\\'{a}ndez et al.\\ (\\cite{martin03}), including optical depth effects, clumpiness, dust extinction, and stellar content. Dust mixed in with the ionized gas has also been proposed as an explanation by Arthur et al.\\ (\\cite{arthur04}, hereafter A04) who argue that significant absorption of ionizing photons by dust grains in the densest \\hii\\ regions flattens the size--density relation.  Previous work on extragalactic \\hii\\ regions suggests that they follow the same size--density correlation as Galactic ones. Kennicutt (\\cite{kennicutt84}) found that \\hii\\ regions in spiral disks follow a $n_\\mathrm{e}\\propto D^{-1}$ relation, but extend the Galactic trend to lower densities and larger sizes. More recently, Gilbert \\& Graham (\\cite{gilbert07}) examined the SSCs in the Antennae galaxies, a prototypical starburst merger, and found that these also follow the $n_\\mathrm{e}\\propto D^{-1}$ relation. However, the SSCs in the Antennae have densities ($\\sim40-400$\\cmcubed) and sizes (25--100\\,pc) that place them on a different location in the $n_\\mathrm{e}$--$D$ plane than the Kennicutt sample. In fact, Gilbert \\& Graham (\\cite{gilbert07}) conclude that the SSCs in the Antennae constitute a new class of massive \\hii\\ regions that is distinct from the Galactic and typical extragalactic population.  In this paper, we examine the size--density relation in extragalactic \\hii\\ regions, and explore several mechanisms which have been proposed to explain it, including clumpiness, dust extinction, and stellar content. In particular, we focus on the \\hii\\ regions in Blue Compact Dwarf galaxies (BCDs); following Gilbert \\& Graham  (\\cite{gilbert07}), we shall refer to such regions as Emission-Line Clusters (ELCs). In fact, most of the current star formation in many BCDs occurs in such ELCs, rather than in the underlying diffuse component. Previously we studied a more limited sample, and found that the \\hii\\ regions in BCDs also follow the same size--density relation namely,  $n_\\mathrm{e}\\propto D^{-1}$ (Hunt et al.\\ \\cite{hunt-cozumel}). Here we triple the sample size used in our previous work, and compare the BCD ELCs with other types of extragalactic \\hii\\ regions, such as those in spiral disks and known SSCs.  Focusing on BCDs enables us to study another aspect of the formation of star clusters. Because they are generally metal poor, with oxygen abundances [$12+\\log (\\mathrm{O/H})$] ranging from 7.2 to 8.5 (Izotov, Thuan, \\& Stasi{\\'n}ska \\cite{izotov07}), BCDs provide a link between high-redshift metal-free primeval galaxies and the metal-enriched SF galaxies in the local universe. Because metallicity is not expected to be the only driver of star-formation and cluster properties, we can compare different metallicities together with other parameters, and better disentangle the effects of metal abundance.  In fact, the process of ionization itself in metal-poor objects is also important in the cosmological context.  Because they are relatively chemically unevolved, \\hii\\ regions in BCDs could enable us to infer some characteristics of the ionization processes at high $z$. Since the typical virial temperature of the first-generation objects in the Universe is $\\la 10^4$ K (e.g., Tegmark et al.\\ \\cite{tegmark97}; Yoshida et al.\\ \\cite{yoshida03}), the photoionization which raises the gas temperature to $\\sim 10^4$ K prohibits the gas from collapsing to form stars (Omukai \\& Nishi \\cite{omukai99}). The gas density structure is modified by the pressure-driven expansion of \\hii\\ regions (Kitayama et al.\\ \\cite{kitayama04}). Such effects on the gas structure are of fundamental importance for considering the subsequent star formation and the reionization of the Universe.  The relation $n_\\mathrm{e}\\propto D^{-1}$ for Galactic and extragalactic \\hii\\ regions can be interpreted as a {\\it constant ionized-gas column density}. However, the physical mechanism that produces the constant column density is not clear, nor is it clear why ELCs in external galaxies and Galactic \\hii\\ regions should follow a similar relation with column density. It is also not understood whether or not there are different classes of \\hii\\ regions which would occupy distinct zones of parameter space in the size--density plane, and if there are, how they might be related.  The aim of this paper is to understand the size--density relation of extragalactic \\hii\\ regions, and place it in the context of star-cluster and ELC formation. The paper is organized as follows. First, in Sect.\\ \\ref{sec:data}, we describe the observational samples of extragalactic \\hii\\ regions, together with the compilation of Galactic \\hii\\ regions for comparison. In Sect.\\ \\ref{sec:overall}, we examine gross trends in the data, which we will later interpret in the light of our evolutionary models. We predict the size--density relation for {\\it static} dusty \\hii\\ regions by assuming a constant luminosity of the central sources in Sect.\\ \\ref{sec:static}. Then, in Sect.\\ \\ref{sec:model}, we extend the model to include pressure-driven expansion and star formation history. Some basic results of this evolutionary model are given in Sect.\\ \\ref{sec:constraint}, and compared with the extragalactic observational data. In Sect.\\ \\ref{sec:discussion} we discuss the models and their implications in various contexts. Finally we give our conclusions in Sect.\\ \\ref{sec:conclusion}. We adopt a Hubble constant of $H_0=73$ km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "\\label{sec:conclusion}  We have investigated the size--density relation of extragalactic \\hii\\ regions, focusing on those in BCDs. Motivated by the similarity of size--density relations of extragalactic \\hii\\ regions with Galactic ones, we have modelled and examined the size--density relation of ionized regions by considering the effects of dust, star formation history, and pressure-driven expansion of \\hii\\ regions. The results have been compared with several samples spanning roughly six orders of magnitude in size and density. We have shown that the entire sample set cannot be understood as an evolutionary sequence with a single initial condition. Rather, the size--density relation reflects a sequence with different initial gas densities. Thus, a hierarchical structure of SF regions with various densities is implied.  We have also found that the size of extragalactic \\hii\\ regions is ``dust-extinction limited'', in the sense that the dust absorption of ionizing photons is significant. This naturally explains the observed size--density relation of \\hii\\ regions as following a constant column density of ionized gas, if the dust-to-gas ratio in \\hii\\ regions is constant. The dust extinction of ionizing photons is particularly severe over the entire lifetime of the compact radio sample with typical densities of $\\ga 10^3$\\cmcubed. This means that the compact radio sample constitutes a different population from the optical samples and that star formation activity in such dense regions would be underestimated or missed entirely if we use the emission from \\hii\\ regions (hydrogen recombination lines, free-free continuum) as the sole indicators of star formation rate."
    },
    "0910/0910.4986.txt": {
        "abstract": "{Deviations from the Bunch-Davies vacuum during an inflationary period can leave a testable imprint on the higher-order correlations of the CMB and large scale structures in the Universe.  The effect is particularly pronounced if the statistical non-Gaussianity is inherently large, such as in models of inflation with a small speed of sound, e.g. DBI.  First reviewing the motivations for a modified vacuum, we calculate the non-Gaussianity for a general action with a small speed of sound.  The shape of its bispectrum is found to most resemble the `orthogonal' or `local' templates depending on the phase of the Bogolyubov parameter. In particular, for DBI models of inflation the bispectrum can have a profound `local' template feature, in contrast to previous results. Determining the projection into the observational templates allows us to derive constraints on the absolute value of the Bogolyubov parameter. In the small sound speed limit, the derived constraints are generally stronger than the constraint obtainable from the power spectrum. The bound on the absolute value of the Bogolyubov parameter ranges from the $10^{-6}$ to the $10^{-3}$ level for $H/\\Lambda_c=10^{-3}$, depending on the specific details of the model, the sound speed and the phase of the Bogolyubov parameter.} ",
        "introduction": "%%%%%%%%%%%%%%%%%%%%%%% Recent years have witnessed a steady increase in the amount of precision data from the cosmic microwave background and large scale structure \\cite{WMAP5}.  This opportunity has motivated the study of both the power spectrum (or two-point correlation) of the primordial density perturbations, essentially probing the free theory of the inflationary Lagrangian, as well as higher-order statistics, probing the interacting theory. The leading statistical probe of interactions is the bispectrum (or three-point correlation, indicating non-Gaussian statistics) and there is currently a large collection of work on the non-Gaussian predictions for specific models of inflation \\cite{ng-singlefield, ng-multifield, ng-curvaton, ng-ekpyrosis, ng-misc}. Interestingly, any detection of non-Gaussianity will immediately rule out all models of single-field slow-roll inflation.  While this would be a monumental discovery in itself, it would also be the starting point for further constraining the inflationary Lagrangian.  Indeed, if the non-Gaussian amplitude is high enough to be observable it opens up a new, and very distinctive, window on the details of inflation \\cite{Komatsu:2009kd}.  In a previous article by two of us (PDM, JPvdS) \\cite{Meerburg} we computed the primordial bispectrum for a single-field slow-roll inflationary model, with the assumption that the initial state is not the Bunch-Davies (BD) vacuum \\cite{Bunch:1978yq}.  This assumption is well-motivated because inflation is by construction an effective theory and therefore there is no `knowledge' of anything beyond a specified cutoff scale in energy.  In expanding backgrounds this energy-cutoff is tantamount to a (possibly momentum-dependent) time-cutoff, allowing one to instead parameterize ignorance as an initial state from which to evolve the dynamics.  As with all effective theories, the details underlying the construction of such a state are irrelevant; the only concern is to use this state to compute the consequences on the bispectrum, and if possible compare theoretical predictions with the data to place constraints on such a deviation.  As first reported in \\cite{Holman2007},  our results in \\cite{Meerburg} confirmed that small initial state modifications can induce surprisingly large corrections to the primordial bispectrum, especially if one includes higher-derivative operators (see also \\cite{ColHol}).  The intuitive explanation is essentially that any deviation from BD introduces particles on sub-horizon scales at the initial stages of inflation, and these particles then have ample time to produce a large non-Gaussian signal in the presence of (irrelevant) interactions.  Unfortunately, this yielded only a rather weak bound of order $|\\beta|<10^{-2}$ on the deviations from the BD state in terms of the Bogolyubov $\\beta$-parameter. This is because although the \\emph{magnitude} of such induced non-Gaussianity was large, its \\emph{shape} in co-moving momentum space was found to substantially depart from standardized observational templates, providing poor overlap with existing constraints. This motivates the need for better templates and possible systematic analysis of oscillatory features in the bispectrum, for which we hope to report some progress in the future.  In this article we generalize the class of models for which initial state modifications are considered.  We opine that initial state modifications are generally natural for these models.  In particular, we will argue that for both branches in DBI inflation, UV \\cite{AST:UVDBI} and IR \\cite{Chen:IRDBI,Chen:2005fe}, small deviations can be expected or at least are worth considering when computing the full bispectrum.  Our starting point will be the most general single-field action, for which the bispectrum has been previously computed in \\cite{SeeLid, Chenetal}.  Specific single-field models (slow-roll, DBI, kinetic) are all limiting cases of this action, and we will be interested in the limit where the speed of sound $c_s$ (that is, the speed at which inflaton perturbations evolve) is significantly reduced relative to $c$, the speed at which metric perturbations evolve. Whereas bispectrum components in standard slow-roll models are typically proportional to $\\epsilon$ and $\\eta$ \\cite{Maldacena}, in the limit of small $c_s$ one finds new contributions to the bispectrum that are proportional to $1/c_s^2$.  Assuming the BD-vacuum these potentially large corrections are typically of the equilateral type and have been previously constrained \\cite{CSZT2007, SSZ2009}.  When one allows for (initially Gaussian) vacuum state modifications, there are yet-additional components to the bispectrum which are also expected to be enhanced in the small speed of sound limit. As a consequence, we will be able to put relatively strong constraints on $|\\beta|$. Importantly, we will also study the effect of an arbitrary phase in the Bogolyubov parameter, which has so far been neglected.  Somewhat surprisingly, we will find that specific choices for the phase can have dramatic effects on the shape of the leading non-Gaussian contribution. As a consequence the phase of the complex Bogolyubov parameter $\\beta$ strongly affects the observational constraints on the absolute value $|\\beta|$ that can be derived. An important restriction in our analysis is that we will have to assume approximate scale-invariance throughout and will not discuss (big or small) departures from scale invariance and how these can be observationally constrained \\cite{Khoury, NG-scaledep}.  This article is organized as follows.  In \\S 2 we will briefly summarize the small-$c_s$ models of inflation and explain how initial state modifications are quite natural in this context.  In \\S 3 we begin from the most general single-field inflationary action and then compute the corrections to the bispectrum.  In \\S 4 we identify the leading-order and sub-leading components and then constrain modifications to the initial state.  Constraints can only be set using already measured bispectra by comparing the theoretical components to observed templates.  In order to achieve this, we compute the leakage via projection of these components into measured components, for which we make use of a formalism first proposed in \\cite{BCZ2004} and optimized in \\cite{Fergusson}.  In \\S 5 we use the two derived constraints on different $f_{\\mathrm{NL}}$ to derive constraints on DBI models with three free parameters.  In \\S 6 we summarize and conclude.  %%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "%%%%%%%%%%%%%%%%%%%%%%%  We have analyzed three-point correlators in single-field inflation with a small speed of sound after introducing a small departure from the standard Bunch-Davies vacuum. Our aim was to derive the best possible constraints on slightly modified vacua in the specific context of small sound speed models like DBI inflation. In previous work \\cite{Meerburg} we concluded that the bispectrum is an excellent probe of vacuum state modifications due to enhancement of the non-Gaussian signal by a factor proportional to (powers of) the ratio of the cutoff scale and the Hubble parameter $p\\equiv \\Lambda_c/H c_s$, which could easily be as large as $10^3$. We have shown that in the small sound speed limit the powers of $p$ that appear can be as large as $p^3$, which adds one factor of $p$ as compared to the $c_s=1$ model. In addition the level of non-Gaussianity in small sound speed models is already large to begin with. Therefore it is not surprising that the constraints that can be derived in such models are in general stronger. The computation of the contributions to the local, equilateral and orthogonal template confirmed this expectation. In general single field models with a small speed of sound the localized enhancement is proportional to $p^3$, and after projecting onto the local and orthogonal template one factor of $p$ is lost, ending up with $p^2$ enhancement. The inclusion of an arbitrary phase in our analysis in the Bogolyubov parameter turns out to be able to reduce the enhancement with one factor of $p$ for very specific values of the phase. For general single field  models with a small speed of sound this implies that for generic phases the constraint on the absolute value of the Bogolyubov parameter equals $|\\beta| \\lesssim 6\\times 10^2 H^2/\\Lambda_c^2$, which is actually independent of the speed of sound. This strong constraint is in fact excluding Bogolyubov parameters linear in $H/\\Lambda_c$ in the small sound speed limit. To be precise, assuming $|\\beta| \\propto \\frac{H^\\star}{\\Lambda_c}$ one obtains a ($c_s$ dependent) {\\it lower} bound on the ratio $\\frac{H}{\\Lambda_c}$, which in the limit of a small sound speed would violate the power spectrum constraint. For the specific value of the phase $\\delta=\\pi/2$, the subleading terms are of the same order and have to be included in the analysis. Similarly, it is the subleading terms that govern the constraint for DBI models, for which the leading term identically vanishes.  In the context of DBI models we performed the same analysis, with a reduced enhancement of $p^2$ in general. After projecting onto the different templates and enforcing the observational bounds this resulted in a constraint on the absolute value of the Bogolyubov parameter $|\\beta| \\lesssim 14 H/(\\Lambda_c c_s)$, which is linear in $H/\\Lambda_c$ and explicitly depends on the speed of sound. This subleading contribution is the one that should also be taken into account for the specific value of the phase $\\delta=\\pi/2$ in the general single field case, since this bound is much stronger than the bound from the second order terms in the leading expression. For general single field inflation with a small speed of sound this results in a very tight bound on allowed values of $|\\beta|$ ranging from $10^{-5}$ to $10^{-3}$. For DBI models of inflation the constraints are only determined by the subleading term as the leading term vanishes. Again, the inclusion of the phase of the Bogolyubov parameter has the effect that for the specific value of $\\delta=0$ the constraint weakens slightly, equal to $|\\beta| \\lesssim 72 H/\\Lambda_c c_s$, but this time remains at the same order in $p$. Overall we conclude that for generic values of $\\delta$ and realistic values of the cutoff this still implies $|\\beta|\\lesssim 10^{-3}-10^{-2}$.  For very specific values of the phase, the bounds could vanish altogether. For these special phases one would have to rely on the power spectrum to set a constraint on $|\\beta|$.  Our main message is that for single field models with a small speed of sound, including DBI, the constraints on vacuum state modifications from the bispectrum for generic values of the phase are already better than the power spectrum constraints, relying on the available templates. The main reason for this is the huge (localized) enhancement that appears after introducing a modified initial state and a corresponding cutoff that regularizes the relevant integrals for the bispectrum calculation. Improved templates, which would have to be sensitive to the oscillatory nature of most of the contributions to the bispectrum, are guaranteed to tighten these constraints significantly. As was previously noted \\cite{Meerburg}, the oscillatory nature of the signal in momentum space suggests that specific, perhaps observable, features could appear in three- and two-dimensional position space. In any case it would be worthwhile to generalize the range of available non-Gaussian shapes that can be compared to the data, including oscillatory signals, on which we hope to report in the future.  Our results confirm that higher-derivative corrections are excellent probes for departures from the standard Bunch-Davies vacuum state. Throughout this paper we assumed that the combination $|k_1 \\eta_0|$ is independent of the actual comoving momenta involved and equal to $\\Lambda_c/H$, in the spirit of the New Physics Hypersurface approach to vacuum state modifications. The reason for this was scale-invariance of the bispectrum, which we relied on to allow for comparison with the available (scale-invariant) template shapes. Fixing $\\eta_0$ instead, as one would do in a Boundary Effective Field Theory approach to vacuum state modifications, immediately results in a scale-dependent bispectrum. It would be interesting to study such scale-dependent scenarios and determine to what extent (future) analysis of 3d large scale structure or 2d CMB data can constrain bispectrum departures from scale-invariance \\cite{NG-scaledep}.  As reported, the bispectrum or three-point function is extremely sensitive to initial vacuum state modifications in the presence of a higher-derivative operator, and there is no reason to think this could not similarly be true for all higher $n$-point functions. A more general perturbative analysis, including higher $n$-point functions \\cite{trispectrum, ChenShiu-etal}, might lead to a hierarchy of (theoretical) constraints on vacuum state modifications, perhaps ruling out departures from the standard Bunch-Davies state altogether. Indeed, for the four-point function (trispectrum) the effect of a vacuum state modification was recently studied in \\cite{ChenShiu-etal}, again showing a strong sensitivity.  We hope that ongoing future work in this direction can help us further understand and identify the phenomenological and theoretical constraints the inflationary vacuum state has to satisfy.  %%%%%%%%%%%%%%%%%%%%%%%"
    },
    "0910/0910.2459_arXiv.txt": {
        "abstract": "We study the emission observed at energies $>$100 MeV of 11 Gamma Ray Bursts (GRBs) detected by the \\fe/Large Area Telescope (LAT) until October 2009. The GeV emission has three main properties: (i) its duration is often longer than the duration of the softer emission detected by the Gamma Burst Monitor (GBM) onboard \\fe\\ [this confirms earlier results from the Energetic Gamma--Ray Experiment Telescope (EGRET)]; (ii) its spectrum is consistent with $F_{\\nu}\\propto \\nu^{-1}$ and does not show strong spectral evolution; (iii) for the brightest bursts, the flux detected by the LAT decays as a power law with a typical slope: $t^{-1.5}$. We argue that the observed $>$0.1 GeV flux can be interpreted as afterglow emission shortly following the start of the prompt phase emission as seen at smaller frequencies. The decay slope is what expected if the fireball emission is produced in the radiative regime, i.e. all dissipated energy is radiated away. We also argue that the detectability in the GeV energy range depends on the bulk Lorentz factor $\\Gamma$ of the bursts, being strongly favoured in the case of large $\\Gamma$. This implies that the fraction of bursts detected at high energies corresponds to the fraction of bursts having the largest $\\Gamma$. The radiative interpretation can help to explain why the observed X--ray and optical afterglow energetics are much smaller than the energetics emitted during the prompt phase, despite the fact that the collision with the external medium should be more efficient than internal shocks in producing the radiation we see. ",
        "introduction": "The {\\it Fermi Gamma Ray Space Telescope (Fermi)} has onboard two instruments: the Large Area Telescope (LAT), sensitive in the 100 MeV -- 100 GeV energy range (and even beyond 100 GeV, for very bright sources, Atwood et al. 2009), and the Gamma Bursts Monitor (GBM), especially designed for the detection of Gamma Ray Bursts (GRBs), sensitive in the 8 keV -- 40 MeV energy range (Meegan et al. 2009). The LAT revealed 12 GRBs above 100 MeV confirming that GRBs can be sources of very high energy photons and that the fraction of GRBs that can be detected at these energies is roughly 10 per cent of those detected by the GBM at lower energies. It was the EGRET instrument, onboard the {\\it Compton Gamma Ray Observatory (CGRO)} the first to detect GRBs above 100 MeV (Fishman \\& Meegan 1995; Kaneko et al. 2008), but it is the much better sensitivity (and reduced dead time) of the LAT to allow us for the first time to try to understand the origin of this emission and to answer the question: does it belong to the prompt phase or is it afterglow emission produced by the fireball colliding with the circum--burst medium? Or has it still another origin?  One of the puzzling features of the high energy emission as revealed by EGRET was that it was long lasting, yet it started during the prompt phase as seen by the Burst Alert and Transient Experiment (BATSE) onboard {\\it CGRO} sensitive in the 30 keV -- 1 MeV energy band. For instance, GRB 940217 emitted $>100$ MeV photons up to 1.5 hours after the prompt phase ended in the BATSE detector. A photon of 18 GeV was received $\\sim$5000 s after the trigger (Hurley et al. 1994), and this was the highest photon energy of a GRB until the \\fe--LAT era. On the other hand, about a third of the high energy photons were received within 120 s, before the end of the prompt phase as detected by BATSE.  Up to now, there have been three LAT--detected GRBs already discussed in the literature. In GRB 080916C (Abdo et al. 2009a), there is evidence that the spectrum from 8 keV to 10 GeV can be described by the same Band function (i.e. two smoothly connected power laws), suggesting that the LAT flux has the same origin of the low energy flux. On the other hand, the flux level of the LAT emission, its spectrum and its long lasting nature match the expectations from a forward shock, leading Kumar \\& Barniol--Duran (2009) to prefer the ``standard afterglow\" interpretation (see also Razzaque, Dermer \\& Finke 2009 for an hadronic model; Zhang \\& Peer 2009 for a magnetically dominated fireball model and Zou et al. 2009 for a synchrotron self--Compton origin).  \\begin{table*} \\caption{ The 12 bursts detected by the \\fe--LAT instrument above 100 MeV, until October 03 2009. Besides their redshifts (when measured) and duration, we give the parameters of the time integrated GBM spectrum collected from the literature and the corresponding reference. Fluences $S$ are in [erg cm$^{-2}$], peak energies $E_{\\rm peak}$ in keV. In column 10 we report the fluence $S$ in the [8 keV--10 MeV] energy range calculated from the spectral parameters of the GBM. Column 11 reports the fluence in the [0.1--100 GeV] energy range obtained from the analysis of the LAT spectra performed in this paper (whose results are given in Tab. 2). We adoped a BAND model for the GBM, and a simple power law of photon slope $\\Gamma$ for the LAT. When $\\beta$ is not indicated, the adopted fitting model is a cut off power--law of photon slope $\\alpha$. $^a$: $S_{\\rm GBM}$ in the [8 keV--30 MeV] energy range $^b$: $S_{\\rm GBM}$ in the [50 --300 keV] energy range; $^c$: $S_{\\rm GBM}$ in the [50 keV--40 MeV] energy range; $^d$: $S_{\\rm GBM}$ in the [50 keV--10 MeV] energy range. The number quoted in the ``Ref.\" column refer to GCN circulars as follows: 8141: van der Horst \\& Connaughton 2008; 8278: van der Horst \\& Goldstein  2008; 8407: Omodei 2008; 8682: Chaplin, van der Horst \\& Preece 2008; 8902: von Kienlin 2009; 9021: Ohno M. et al. 2009; 9057: Rau, Connaughton \\& Briggs 2009; 9336: Guiriec, Connaughton \\& Briggs 2009; 9579: von Kienlin 2009; 9866: Bissaldi \\& Connaughton 2009; 9933: Bissaldi 2009; 9983: Rau 2009. } \\label{tab0} \\centering \\begin{tabular}{l l l l l l l l l l l} \\hline \\hline GRB     &$z$     &$T_{90}$ &$S_{\\rm GBM}$ & $\\alpha_{\\rm GBM}$ &$\\beta_{GBM}$ &$E_{\\rm peak}$   &Ref   &$E_{\\rm \\gamma, iso}$ &$S_{\\rm GBM}$  &$S_{\\rm LAT}$ \\\\ &        &s        &                        &                  &                &keV           &     &erg    &8--10$^4$ keV  &0.1--100 GeV \\\\ \\hline 080825C &...   &22         &2.4e--5                 &--0.39$\\pm$0.04   &--2.34$\\pm$0.09 &155$\\pm$5     &8141 &...    &(3.4$\\pm$0.3)e--5   &(9.5$\\pm$4)e--6  \\\\ 080916C &4.35  &66         &1.9e--4$^a$             &--0.91$\\pm$0.02   &--2.08$\\pm$0.06 &424$\\pm$24    &8278 &5.6e54 &(1.6$\\pm$0.2)e--4   &(7$\\pm$1)e--5  \\\\ 081024B &...   &0.8        &(3.4$\\pm$0.1)e--7       &--0.70$\\pm$0.13   &...             &1583$\\pm$520  &8407 &...    &(3.2$\\pm$0.1)e--6   &(3$\\pm$2)e--6\\\\ 081215  &...   &$\\sim$90   &(2.8$\\pm$0.5)e--6$^b$   &--0.14$\\pm$0.26   &...             &139$\\pm$14    &8682 &...    &                    &... \\\\ 090217  &...   &32.8       &(3.08$\\pm$0.03)e--05    &--0.845$\\pm$0.023 &--2.86$\\pm$0.3  &610$\\pm$32    &8902 &...    &(3.8$\\pm$0.4)e--5   &(4.2$\\pm$1.6)e--6\\\\ 090323  &3.57  &$\\sim$150  &(1.00$\\pm$0.01)e--4     &--0.89$\\pm$0.03   &...             &697$\\pm$51    &9021 &3.4e54 &(1.32$\\pm$0.03)e--4 &(3.6$\\pm$0.8)e--5 \\\\ 090328  &0.736 &$\\sim$25   &(8.09$\\pm$0.10)e--5     &--0.93$\\pm$0.02   &--2.2$\\pm$0.1   &653$\\pm$45    &9057 &2.1e53 &(1.52$\\pm$0.02)e--4 &(3.3$\\pm$2)e--5 \\\\ 090510  &0.903 &1          &(3.0$\\pm$0.2)e--5$^c$   &--0.80$\\pm$0.03   &--2.6$\\pm$0.3   &4400$\\pm$400  &9336 &5.0e52 &(2.3$\\pm$0.2)e--5   &(3.7$\\pm$0.7)e--5\\\\ 090626  &...   &70         &(3.5$\\pm$0.1)e--5       &--1.2$\\pm$0.02    &--1.98$\\pm$0.02 &175$\\pm$12    &9579 &...    &(6.0$\\pm$0.2)e--5   &(9.6$\\pm$6)e--6\\\\ 090902B &1.822 &$\\sim$21   &(3.74$\\pm$0.03)e--4$^d$ &--0.696$\\pm$0.012 &--3.85$\\pm$0.25 &775$\\pm$11    &9866 &4.4e54 &(5.4$\\pm$0.04)e--4  &(5.9$\\pm$0.6)e--4\\\\ 090926A &2.106 &20$\\pm$2   &(1.45$\\pm$0.04)e--4     &--0.75$\\pm$0.01   &--2.59$\\pm$0.05 &314$\\pm$4     &9933 &2e54   &(1.9$\\pm$0.05)e--4  &(4.3$\\pm$0.8)e--5\\\\ 091003  &0.897 &21$\\pm$0.5 & (3.76$\\pm$0.04)e--5    &--1.13$\\pm$0.01   &--2.64$\\pm$0.24 &86.2$\\pm$23.6 &9983 &8.7e52 &(4.16$\\pm$0.03)e--5 &(1.3$\\pm$0.8)e--5 \\\\ \\hline \\hline \\end{tabular} \\end{table*}  \\begin{figure*} \\hskip 0.3 cm \\psfig{figure=plot_lc_new.ps,width=18cm} % \\vskip -0.2 cm \\caption{ Light curves of the 11 GRBs detected by LAT plus GRB 940217, as detected by EGRET (bottom right panel). The hatched region represents the duration ($T_{90}$) of the emission detected by the GBM in the 8 keV--40 MeV energy range (for GRB 940217 it refers to the emission detected by BATSE). Times are in the observer frame for all bursts and arrows represent 2$\\sigma$ upper limits.} \\label{lc} \\end{figure*}   In the short bursts GRB 090510 the spectrum in the LAT energy range is not the extrapolation of the flux from lower energies, but is harder, leading  Abdo et al. (2009b) to propose a synchrotron self--Compton interpretation for its origin. Instead we (Ghirlanda, Ghisellini \\& Nava 2009) proposed that the LAT flux is afterglow synchrotron emission, on the basis of its time profile and spectrum (see also Gao et al. 2009; De Pasquale et al. 2009).  Finally, the LAT flux of GRB 090902B decays as $t^{-1.5}$ (Abdo et al. 2009c), it lasts longer than the flux detected by the GBM, and its spectrum is harder than the extrapolation from lower frequencies, making it a good candidate for an afterglow interpretation, despite the arguments against put forward by Abdo et al. (2009c), that we will discuss in this paper. Moreover, in GRB 090902B there is evidence of a soft excess (observed in the GBM spectrum below 50 keV) which is spectrally consistent with the extrapolation at these energies of the LAT spectrum.  As the few examples above demonstrate, there is no consensus yet on the nature of the high energy emission of GRBs. Since only three of the nearly dozen bursts detected by the LAT have already been discussed in the literature, we present here a study of the entire sample of bursts detected at high energies by the LAT. We will construct the light curves of the high energy flux and the spectral shape in the 0.1--100 GeV energy range, to find if there are properties that are common among different bursts that can help to understand their nature.  Indeed, we believe that a consistent scenario emerges: the LAT spectra are often inconsistent with the extrapolation of the GBM spectra (except two cases) and the light curves can be often described by a power law decay in time, i.e. $F_{\\rm LAT}\\propto t^{-\\alpha}$, with a slope often close to $\\alpha=1.5$. In the brightest cases also the rising part is visible, and is consistent with $F_{\\rm LAT}\\propto t^2$. These are, in our opinion, strong indications of the afterglow nature of the LAT emission. Furthermore, we suggest that GRBs with a flux decaying as $F_{\\rm LAT}\\propto t^{-1.5}$, and with a spectral slope around unity [i.e. $F(\\nu)\\propto \\nu^{-1}$], are emitting in the radiative regime of a forward shock. We will also point out the role that the electron--positron pair production process has in establishing the radiative regime. Finally, we will discuss the consequences of our findings.  We adopt a cosmology with $h=\\Omega_\\Lambda=0.7$ and $\\Omega_{\\rm M}=0.3$ and the convention $Q = 10^x Q_x$, using cgs units. ",
        "conclusions": "The found bulk Lorentz factors are in the range 630--900 for the long bursts, and 2000 for the short GRB 090510. We believe that these relatively large values are the key to understand why only a minority of bursts are detectable by the LAT. A large bulk Lorentz factor, in fact, means an early peak time of the afterglow (see Eq. \\ref{liso3} and Eq. \\ref{liso4}), and this in turn means a large flux. {\\it Faster fireballs have brighter afterglows}. This is true for adiabatic as well as radiative fireballs. If the emission occurs in the radiative regime then the afterglow will be brighter still, since all the energy dissipated in the external shock is radiated away.  If the circum--bursts medium is enriched by electron--positron pairs, we have a more favourable set up for a radiative process. If the acceleration mechanism divides its energy to all particles, then leptons should receive a total energy exceeding the one given to protons. But this may be only one of the means to have a radiative fireball. An alternative is to have a strong coupling between electrons and protons, with an efficient energy flow from protons to electrons. In any case, we can easily test if pairs are indeed important by simply comparing the general properties of the early afterglow for bursts of different $E_{\\rm peak}$ and high energy index $\\beta$, since only those bursts whose prompt phase photon energies exceed $m_{\\rm e}c^2$ should efficiently populate the circumburst medium by pairs. As an example, we may test if the high energy emission is present only in GRBs of high $E_{\\rm peak}$ (in the rest frame) as it appears to be the case until now, or if it occurs also for bursts with a small $E_{\\rm peak}$. If this will occur, and if the flux will decay with a slower rate than $t^{-10/7}$, then we will have an indication of a fast fireball that emits adiabatically because of no pairs--enrichment of the circum--bursts medium. In other words, a possible test of the idea of having radiative afterglows because of pair enrichment is to find a different time decay for the high energy emission in classical GRBs whose prompt phase emission extends to high energies and X--ray flashes, characterised by relatively small values of $E_{\\rm peak}$.   The radiative interpretation could ease the efficiency problem of the afterglow phase. This problem concerns the ratio of the energetics emitted during the prompt and afterglow phases, that is much larger than unity (e.g. Zhang et al. 2007). According to the standard internal/external shock scenario one expects the opposite, since external shocks should be much more efficient than internal ones to dissipate the kinetic energy of the fireball. These estimates were based on the observed X--ray afterglow energetics (see e.g. Willingale et al. 2007; Ghisellini et al. 2009), and we can now revise them including the much more powerful high energy $\\gamma$--ray emission, bringing the total afterglow energetics to be roughly equal to the prompt phase one. Furthermore, if the fireball is indeed radiative in the first phases, with a consequent fast decay, we can understand why the afterglow emission at later times and at other frequencies is so faint.  According to our findings, bursts detected by the LAT may be the ones with the largest $\\Gamma$, and can be used to explore the high--end $\\Gamma$--distribution. On the other hand, one can wonder about the possibility to detect with the LAT bursts with relatively smaller $\\Gamma$, smaller high energy luminosities and with light curves peaking at larger peak times. Even if rare, nearby objects with these properties might be still detectable, offering a direct way to test our ideas: even if they should be characterised by much lower peak luminosities in the LAT, they should have LAT/GBM fluence ratios similar to those presented in this paper, and lower values of $\\Gamma$.  One of the argument put forward against the afterglow interpretation of the high energy flux is its variability, that according to Abdo et al. (2009c) can have a timescales $\\Delta t_{\\rm var}$ as short as 90 ms. If true, this is certainly a severe problem for the afterglow interpretation. On the other hand the knowledge of $\\Delta t_{\\rm var}$ is limited by the few number of received photons. When the entire light--curve, lasting for a few hundreds seconds, is composed by a few hundreds events, one can define a very short $\\Delta t_{\\rm var}$ only if there is an exceptional ``bunching\" of photons in contiguous time--bins, and we do not see it in the bursts we analysed.   Finally, we would like to emphasise the importance of establishing, in general, if the high and low energy emission are produced  by the same electrons at the same time or instead if they are produced by different electrons at different times. As the study of GRB 090510 (Abdo et al. 2009b; Ghirlanda et al. 2009) has demonstrated, we are reaching the required data quality to put strong constraints on the theories predicting the violation of the Lorentz invariance at small scales, that can be tested by comparing the possible delay of the {\\it arrival} times of high energy photons. The critical issue about these studies is to know exactly the {\\it generation} time of the high with respect to low energy emission. Therefore it becomes crucial to establish if the flux received by the LAT is the extension in energy of the prompt phase emission or if it is afterglow radiation."
    },
    "0910/0910.5475_arXiv.txt": {
        "abstract": "We present an updated catalogue of M31 globular clusters (GCs) based on images from the Wide Field CAMera (WFCAM) on the UK Infrared Telescope and from the Sloan Digital Sky Survey (SDSS). Our catalogue includes new, self-consistent \\textit{ugriz} and K-band photometry of these clusters. We discuss the difficulty of obtaining accurate photometry of clusters projected against M31 due to small scale background structure in the galaxy. We consider the effect of this on the accuracy of our photometry and provide realistic photometric error estimates. We investigate possible contamination in the current M31 GC catalogues using the excellent spatial resolution of these WFCAM images combined with the SDSS multicolour photometry. We identify a large population of clusters with very blue colours. Most of these have recently been proposed by other work as young clusters. We distinguish between these, and old clusters, in the final classifications. Our final catalogue includes 416 old clusters, 156 young clusters and 373 candidate clusters. We also investigate the structure of M31's old GCs using previously published King model fits to these WFCAM images. We demonstrate that the structure and colours of M31's old GC system are similar to those of the Milky Way. One GC (B383) is found to be significantly brighter in previous observations than observed here. We investigate all of the previous photometry of this GC and suggest that this variability appears to be genuine and short lived. We propose that the large increase in its luminosity my have been due to a classical nova in the GC at the time of the previous observations in 1989. ",
        "introduction": "Globular clusters (GCs) are among the oldest known stellar systems. They typically have ages similar to those of their host galaxies, making them ideal probes into galaxy formation and evolution. The properties of GCs vary significantly. However, individual clusters contain populations of stars with similar ages and metallicities. This makes them unique locations for studying stellar evolution.  The Milky Way's GCs still represent the best studied GC system. While the study of these clusters has led to many advances, the Milky Way contains relatively few GCs \\citep[$\\sim$150 GCs:][]{Harris96}, many of which have high foreground extinction, making them hard to study. By determining the properties of extragalactic GCs, we are able to study a more diverse population and ensure our current conclusions are not biased by the Milky Way's clusters being atypical.  For extragalactic GCs, it is very difficult to resolve individual stars in the clusters. However, it is possible to estimate many properties of a GC from its integrated light. For example: the masses of GCs can be estimated by assuming a mass to light ratio; combined optical and near infrared colours of GCs can be used to (at least partially) break the age and metallicity degeneracy and estimate these parameters \\citep[e.g.][]{Puzia02,Jiang03,Hempel07}; and their structural parameters can be estimated by fitting their density profiles \\citep[e.g.][]{Barmby07,Jordan07,McLaughlin,Peacock09}. The colours of GCs and GC candidates are also very useful in selecting genuine GCs from stellar asterisms and background galaxies. Good multi-wavelength photometry of GCs is therefore highly desireable.   \\subsection{The M31 GC system}  The proximity of M31, and its relatively large GC population compared with the Milky Way \\citep[$\\sim$400$:$][]{Barmby00}, makes it the ideal location to study extragalactic globular clusters. Its clusters have been the focus of many studies dating back to the early work of \\citet{Hubble32} and \\citet{Vetesnik62}. However, attempts to study its clusters have faced several challenges. Photometry of these clusters is complicated by many of them being projected against the bright and non-uniform structure of M31 itself. The galaxy's proximity also results in the GC system extending over a wide region of the sky, with clusters recently found beyond 4$^{\\circ}$ from the centre of the galaxy \\citep{Huxor08}. This means that surveys with large fields of view are required in order to study the GC system. It is also difficult to confirm GCs in M31 based on spectroscopy alone. This is because the velocity distribution of its GC system overlaps that of Milky Way halo stars.  Over the past decades there have been several large catalogues of M31's GCs including those of: \\citet{Battistini87,Barmby00,Galleti04,Kim07}. In addition to these catalogues many new clusters and candidates have been proposed \\citep[e.g.][]{Battistini93,Mochejska98,Barmby02,Galleti06,Galleti07,Huxor08,Caldwell09}. These studies have made considerable progress in removing contamination from the cluster catalogues due to either background galaxies \\citep[e.g.][]{Racine91,Barmby00,Perrett02,Galleti04,Kim07,Caldwell09} or stars and asterisms both in the Milky Way and M31 itself \\citep[e.g.][]{Barmby00,Cohen05,Huxor08,Caldwell09}. However, despite this work, it is likely that there remains significant contamination in the current catalogues of M31 clusters, especially at the faint end of the GC luminosity function. These studies have also resulted in a large number of unconfirmed candidate clusters \\citep[currently over 1000:][]{Galleti04} whose true nature remains uncertain.  It has been known for many years that some of the proposed GCs in M31 have very blue colours. Recent work has identified that a large number of the clusters in the current catalogues are young clusters \\citep[e.g.][]{Beasley04,Fusi_Pecci05,Rey07}. A comprehensive catalogue of young clusters in M31 has recently been published from the spectroscopic survey of \\citet{Caldwell09}. Compared with the young open clusters in the Milky Way, these clusters have relatively high masses ($<10^{5}M_{\\sun}$), akin to the young clusters observed in the Large Magellanic Cloud. A recent \\textit{Hubble Space Telescope (HST)} study of 23 of these young clusters suggested that on average they are larger and more concentrated than typical old clusters \\citep{Barmby09}. Most of the clusters studied by \\citet{Barmby09} are found to have dissolution timescales of less than a few Gyr and are therefore not expected to evolve into typical old globular clusters. Whether these clusters are massive open clusters, young globular clusters or a mix of both, it is clear that they represent a different population to the classical old GCs also observed in M31 (which are the focus of this study). We therefore distinguish between the two populations in our classifications and conclusions.  \\begin{figure*} \\includegraphics[width=170mm,angle=0]{m31_coverage_d25.eps} \\caption{Coverage of the SDSS and WFCAM images used. For reference all objects listed in the RBC are shown in red. Only images covering the locations of these objects were extracted from the SDSS archive. The green ellipse indicates the D$_{25}$ ellipse of M31. The grid represents $2^{\\circ}\\times2^{\\circ}$ squares on the sky ($13.6\\times13.6$~kpc at the distance of M31) and highlights the spatial extent of the GC system. } \\label{fig:images} \\end{figure*}  While these previous studies have provided a wealth of information on the M31 GC system they have also resulted in a rather heterogeneous sample. For example, the excellent and commonly used Revised Bologna Catalogue (hereafter RBC) of M31 GCs by \\citet{Galleti04} includes photometry from many different authors using different telescopes and in some cases different (homogenised) filters. This has been previously noted by \\citet{Caldwell09} (hereafter C09) who recently published new V-band photometry for many RBC sources which were located in the Local Group Galaxy Survey images of M31 \\citep[LGGS:][]{Massey06}. While this work provides excellent deep V-band photometry, the survey does not cover the outer clusters and candidates and does not provide colour information. The most complete set of optical colours, derived in a consistent manner, are still that of the Barmby catalogue \\citep{Barmby00}. This work presented self-consistent UBVRI colours for many of their clusters. However, it is incomplete in some of these bands and only provides new photometry 285 clusters. For these reasons we chose to produce new, self-consistent, optical photometry for the proposed GCs and GC candidates in the RBC using images from the Sloan Digital Sky Survey (SDSS). The excellent calibration and large field of view of this survey is ideal for studying such an extended system. Details of this photometry are presented in section 2.2.  The study of M31's GCs in the near infrared (NIR) is very useful both for confirming genuine GCs and for estimating their ages and metallicities \\citep[e.g.][]{Barmby00,Galleti04,Fan06}. The first major survey of M31's GCs in the NIR was by \\citet{Barmby00} who used pointed observations of individual clusters to obtain K-band photometry of 228 of their clusters. More recently \\citet{Galleti04} obtained NIR photometry in the J,H and K-bands of 279 of their confirmed GCs from the 2 Micron All Sky Survey (2MASS). The spatial coverage of 2MASS makes it ideal for such a project. However, the survey is relatively shallow and has relatively poor spatial resolution. We have obtained new deep K-band photometry using the Wide Field Camera on the UK Infrared Telescope to determine the K-band magnitude of M31's GCs across the entire GC luminosity function. Some results of this survey are already published in \\citet{Peacock09}. In addition to providing the first K-band photometry for 126 GCs marked as confirmed in the RBC, the excellent spatial resolution of these images is very useful for removing stellar sources from genuine clusters, and for investigating the density profiles of the clusters. Details of this new K-band photometry are presented in section 2.3, while the classifications of the proposed clusters and candidates are considered in section 3.  The proximity of M31 makes it the ideal location for studying the structure of extragalactic GCs. Determination of the size and density of GCs is very useful in investigating: stellar evolution; galaxy formation and evolution; constraining N-body simulations; and investigating exotic objects in GCs like X-ray binaries, blue stragglers and horizontal branch/extreme horizontal branch stars. The structural parameters for some of M31's GCs have been measured using \\textit{Hubble Space Telescope (HST)}: Faint Object Camera (FOC) images of 13 clusters \\citep{Fusi_Pecci94}; Wide Field Planetary Camera (WFPC2) images of 50 clusters \\citep{Barmby02}; and Advanced Camera for Surveys (ACS) and Space Telescope Imaging Spectrograph (STIS) images of 15 and 19 clusters respectively \\citep{Barmby07}. In \\citet{Peacock09} we presented the results of fitting PSF convolved King models \\citep{King66} to the ground based WFCAM images used here in order to estimate the structural parameters for 239 clusters. In section 4 we discuss the structure of M31's GCs. ",
        "conclusions": "Our final catalogue includes 416 old GCs. Where detected, we provide self consistent \\textit{ugriz} and K-band photometry for the proposed clusters and candidate clusters. We note the difficulty in providing accurate photometry for some of these clusters due to the complex background of M31. We highlight the need for good spatial resolution in order to remove contamination from non cluster light when obtaining integrated magnitudes. Where available, we find our photometry to be consistent with that previously published. From our multicolour photometry, we confirm the population of very blue clusters identified previously. We show that these colours are consistent with their spectroscopic classification by C09 as young clusters. We note that many of these clusters look like resolved asterisms in our K-band images. However, some of these are confirmed by HST images to be genuine clusters. Higher spatial resolution optical images than available here are required in order to confirm their nature as genuine young clusters.  We have identified that many of the confirmed clusters from \\citet{Kim07} are likely stellar sources (we retain only 27 of their 111 confirmed clusters as old clusters). We also identify 10 confirmed clusters in the RBC as likely stellar sources. While we have considered the classifications from K07 and the RBC separately in this paper, we caution that \\textit{all} of the objects confirmed by K07 to be clusters are included as confirmed clusters in the current version of the RBC (v3.5). We also provide new classifications for many of the cluster candidates proposed by this previous work. We identify many of these candidates to be stars and reduce the number of unclassified candidate clusters to 357.  Taking extinction and selection effects into account, we find both the colours and structure of the old M31 GC system are consistent with the Milky Way's. We note a potential lack of both core collapsed and very extended GCs in our M31 sample. We caution that some (or all) of this effect may be due to selection effects in identifying these clusters, or difficulties in accurately measuring their parameters, rather than an intrinsic difference in the populations."
    },
    "0910/0910.0857_arXiv.txt": {
        "abstract": "Radio AGN feedback in X-ray cool cores has been proposed as a crucial ingredient in the evolution of baryonic structures. However, it has long been known that strong radio AGNs also exist in ``noncool core'' clusters, which brings up the question whether an X-ray cool core is always required for radio feedback. We present a systematic analysis of 152 groups and clusters to show that every BCG with a strong radio AGN has an X-ray cool core. Those strong radio AGNs in the center of the ``noncool core'' systems identified before are in fact associated with small X-ray cool cores with typical radii of $<$ 5 kpc (we call them coronae). Small coronae are most likely of ISM origin and they carry enough fuel to power radio AGNs. Our results suggest that the traditional cool core / noncool core dichotomy is too simple. A better alternative is the cool core distribution function with the enclosed X-ray luminosity. Other implications of our results are also discussed, including a warning on the simple extrapolation of the density profile to derive Bondi accretion rate. ",
        "introduction": "The importance of AGN outflows for cosmic structure formation and evolution has been widely appreciated recently. AGN outflows may simultaneously explain the antihierarchical quenching of star formation in massive galaxies, the exponential cut-off at the bright end of the galaxy luminosity function, the $M_{\\rm SMBH} - M_{\\rm bulge}$ relation and the quenching of cooling-flows in cluster cores (e.g., \\citep{cro06}; \\citep{mn07}), as also widely discussed in this conference. It has been suggested that FRI and low-excitation FRII radio galaxies are fueled through accretion of hot gas in the so-called ``radio mode'' \\citep{chu05}\\citep{best05}\\citep{cro06}\\citep{allen06}\\citep{hard07}. This naturally brings out the following observational questions: {\\em What fuels the strong radio AGN in groups and clusters, especially those not in large cool cores? Is an X-ray cool core always required for strong radio AGN in groups and clusters?}  \\begin{figure} \\includegraphics[height=.21\\textheight]{a3391_all1.eps} \\caption{One example of strong radio AGNs in the center of a cluster without a large X-ray cool core (A3391). \\chandra\\ image around the BCG is shown on the left, while the middle panel shows the radio AGN (the small box shows the field of the left panel). The radio AGN is 2.5 times more luminous than Perseus's at 1.4 GHz, but the ICM cooling time at $r >$ 4 kpc is 17 Gyr. The X-ray point-like source centered on the BCG is in fact a small cool core, revealed by its spectrum shown on the right (clearly a thermal origin). As shown by \\citep{sun09}, the small cool core of A3391 has enough gas to fuel its radio AGN, although its total luminosity is only 0.062\\% of Perseus's central cool core. } \\end{figure}  First of all, it has long been known that there are many strong radio AGNs in the center of ``noncool core'' clusters, or clusters without a large cool core. Some examples are in the HIFLUGCS sample\\citep{Mit09}, e.g., A3391, A3395s, A2634 and A400. \\citep{Mit09} cited cluster merger and other mechanisms to explain these ``outliers''. One example is shown in Fig. 1 (A3391). There is an X-ray point-like source at the position of the BCG. Is it an X-ray AGN? A spectral analysis clearly reveals a significant iron-L hump, implying a thermal origin. Its properties are very similar to many other coronae (or mini-cooling cores) discussed in \\citep{sun07} (also see the references there). Thus, we have a case in sharp contrast to the Perseus cluster, a much much smaller X-ray cool core associated with a stronger radio AGN. Is this case unique? Not at all! \\citep{sun09} presents many more similar cases from a sample study of 152 nearby groups and clusters from the \\chandra\\ archive. As shown in the left panel of Fig. 2, all 69 BCGs with a strong radio AGN ($L_{\\rm 1.4 GHz} > 2 \\times 10^{23}$ W Hz$^{-1}$, no high-excitation FRII galaxies) have X-ray cool cores with a central isochoric cooling time of $<$ 1 Gyr \\citep{sun09}. This conclusion also holds in the B55 and the extended HIFLUGCS samples \\citep{sun09}. The BCG cool cores can be divided into two classes, large ($r_{\\rm 4 Gyr} >$ 30 kpc) and luminous ($L_{\\rm 0.5-2 keV} \\geq 10^{42}$ ergs s$^{-1}$) cool cores like Perseus's cool core, or small ($\\leq$ 4 kpc in radius typically) coronae like A3391's. We call them the large-cool-core (LCC) class (the focus of this conference) and the corona class. The gas of the former class is primary of ICM origin, while the latter one is of ISM origin. More detail on these two classes can be found in \\citep{sun09}.  \\begin{figure} \\includegraphics[height=.38\\textheight]{f2.eps} \\caption{{\\bf Left}: The upper panel shows the rest-frame 0.5-2 keV luminosity of the cool core (within a radius where the cooling time is 4 Gyr) of the BCG vs. the 1.4 GHz luminosity of the BCG. Red filled circles are for $kT>$4 keV clusters. Blue triangles are for $kT$=2-4 keV poor clusters. Green stars are for $kT<$2 keV groups. Crosses represent upper limits in both axes. The lower panel shows the histogram for all BCGs (upper limits included) with two classes marked, while the histogram in red is for $L_{\\rm 1.4GHz} > 10^{24}$ W Hz$^{-1}$ BCGs. This histogram can be regarded as a raw {\\bf cool-core distribution function}. There are two classes of BCG cool cores (shown in orange ellipses): the large-cool-core (LCC) class and the corona class. Above $L_{\\rm 1.4GHz}$ of 2$\\times10^{23}$ W Hz$^{-1}$, every BCG has a confirmed cool core, either in the LCC class or in the corona class (see \\citep{sun09} for detail). Radio AGNs in the corona class can be as luminous as those in LCC class. {\\bf Right}: The electron density of ESO~137-006's corona at the Bondi radius (62 pc) vs. the slope of the density profile in the innermost bin ($n_{\\rm e} \\propto r^{-\\alpha}$). The solid line is the relation constrained from the observed normalization of the innermost bin. The dashed line is for the relation when the normalization is doubled, which should over-estimate the PSF correction. The dotted line is from the boundary condition of the density in the second innermost bin (see Fig. 7 of \\citep{sun09}). The intersections of these lines imply the density at the Bondi radius as $\\sim$ 1.5 cm$^{-3}$. } \\end{figure} ",
        "conclusions": ""
    },
    "0910/0910.0250_arXiv.txt": {
        "abstract": "We examine the impact of gas pressure on the transverse coherence of high-redshift ($2\\leq z \\leq 4$) \\lyman\\ forest absorption along neighboring lines of sight that probe the gas Jeans scale (projected separation $\\Delta r_p\\leq 500 h^{-1}$\\,kpc comoving; angular separation $\\Delta\\theta\\lesssim 30\\arcsec$).  We compare predictions from two smoothed particle hydrodynamics (SPH) simulations that have different photoionization heating rates and thus different temperature-density relations in the intergalactic medium (IGM). We also compare spectra computed from the gas distributions to those computed from the pressureless dark matter.  The coherence along neighboring sightlines is markedly higher for the hotter, higher pressure simulation, and lower for the dark matter spectra.  We quantify this coherence using the flux cross-correlation function and the conditional distribution of flux decrements as a function of transverse and line-of-sight (velocity) separation.  Sightlines separated by $\\Delta\\theta\\lesssim 15\\arcsec$ are ideal for probing this transverse coherence.  Higher pressure decreases the redshift-space anisotropy of the flux correlation function, while higher thermal broadening increases the anisotropy.  In contrast to the longitudinal (line-of-sight) structure of the \\lya\\ forest, the transverse structure on these scales is dominated by pressure effects rather than thermal broadening.  With the rapid recent growth in the number of known close quasar pairs, paired line-of-sight observations offer a promising new route to probe the IGM temperature-density relation and test the unexpectedly high temperatures that have been inferred from single sightline analyses. ",
        "introduction": "\\label{sec:intro} The $1.5\\lesssim z\\lesssim 6$ intergalactic medium (IGM) is most commonly studied via the \\lyman\\ forest, which arises from \\lya\\ absorption of neutral hydrogen along the line of sight to some distant source (e.g., a quasar).  Most of the information we have on the nature of the IGM is from sightlines piercing physically distinct regions of the IGM.  Absorption features have finite widths, but from individual sightlines it is difficult to separate the contribution of bulk velocities (Hubble flow and peculiar velocities) from those of thermal broadening.  Close quasar pairs can break this degeneracy by probing the transverse structure of the IGM.  In the mid-1990s, several studies of quasar groups and lensed quasars definitively showed that the absorbing structures are coherent over hundreds of kpc \\citep{bechtold94,dinshaw94,fang96,charlton97,crotts98,dodorico98}.  These observations provided critical support for the physical picture of the \\lya\\ forest then emerging from cosmological simulations \\citep{cen94,zhang95,hernquist96} and associated analytic descriptions \\citep{rauch95,reisenegger95,bi97,hui97a}, in which most \\lya\\ absorption arises in a continuously fluctuating medium of low density gas rather than in a system of discrete clouds.  More recently, observations of pairs have been suggested as ways of investigating measuring the cosmological constant via the \\citet{alcock79} effect \\citep{mcdonald99}, the matter power spectrum on small scales \\citep{viel02a}, and the IGM temperature-density relation.  It is on the last of these that we focus in this paper.  Theoretical models predict that the low density IGM should have a power-law ``equation of state,'' \\begin{equation}\\label{eqn:rhot} T=T_0(1+\\delta)^{\\alpha}, \\end{equation} with denser gas being hotter than less dense gas \\citep[$\\alpha > 0$,][]{katz96,miralda96,hui97b}.  Although this temperature-density (\\rhot) relation is difficult to measure, multiple observations suggest that it has a higher normalization and a shallower slope than that expected using the most straightforward assumptions about photoionization heating \\citep{schaye00,mcdonald01,theuns02,bolton08}. Because the degree of small-scale transverse coherence is set by the Jeans length, and the Jeans length depends on the temperature of the gas, studying the transverse structure of the \\lya\\ forest might give insight into the IGM \\rhot\\ relation (J.\\ Hennawi, private communication, 2007).  In Peeples et al.\\ (2009, hereafter Paper~I), we show that while thermal broadening and pressure support both affect the longitudinal structure of the \\lya\\ forest, thermal broadening dominates.  In this paper we investigate the effects of the temperature-density relation via pressure support and thermal broadening on the transverse small-scale structure of the \\lya\\ forest.  The gas temperature affects the Jeans length $\\lambda_J$ (and the comoving Jeans length $\\lambda_{J,{\\rm comv}}$) via \\begin{eqnarray} \\lambda_J &\\equiv& c_s\\sqrt{\\frac{\\pi}{G\\rho}}\\label{eqn:jeans}\\\\ \\Rightarrow\\lambda_{J,{\\rm comv}} & = & (1+z)\\sigma_{\\rm th}H_0^{-1}\\sqrt{\\frac{5\\pi}{3}} \\left[\\frac{3}{8\\pi}\\Omega_{m,0}(1+z)^3(1+\\delta)\\right]^{-1/2}\\nonumber\\\\ &= & 782\\,h^{-1}\\,\\mbox{kpc}\\label{eqn:jeanscomove}\\\\ && \\times\\left(\\frac{\\sigma_{\\rm th}}{11.8\\,\\rm{km}\\,\\rm{s}^{-1}}\\right)\\left[\\left(\\frac{\\Omega_{m,0}(1+\\delta)}{0.25 \\times(1+0)}\\right)\\left(\\frac{1+z}{1+3}\\right)\\right]^{-1/2}, \\nonumber \\end{eqnarray} where $1+\\delta\\equiv\\rho_{\\rm gas}/\\bar{\\rho}_b$ is the gas overdensity and $c_s=\\sqrt{[5kT]/[3m]}=\\sigma_{\\rm th}\\sqrt{5/3}$ is the speed of sound in an ideal gas expressed as a multiple of the 1-D thermal velocity $\\sigma_{\\rm th}$, which we have normalized to correspond to $10^4$\\,K \\citep{miralda96,schaye01,desjacques05}.  While thermal broadening affects the observed IGM by smoothing the \\lya\\ forest in one-dimension (namely, along the line of sight), pressure support smooths the physical gas distribution in all three dimensions.  Therefore, while we found in Paper~I that $\\sigma_{\\rm th}$ dominates the longitudinal \\lya\\ forest structure, we expect $\\lambda_J$ to dominate the transverse structure.  Our simulations indicate that the ``effective'' Jeans length in the \\lya\\ forest is smaller than that given in Equation~(\\ref{eqn:jeanscomove}) by a factor of a few, probably owing to a combination of geometric factors, the universe expanding on the same timescale as the gas evolves, and the contribution of dark matter to the gravitational forces \\citep[see also][]{gnedin98}.  For $\\Omega_m=0.25$ and $\\Omega_{\\Lambda}=0.75$, the relation between angular separation $\\Delta\\theta$ and comoving transverse separation $R$ at $z=2$--4 is approximately \\begin{equation}\\label{eqn:dtheta} \\Delta\\theta\\approx 4.4\\arcsec\\left(\\frac{1+z}{4}\\right)^{0.6}\\times\\left(\\frac{R}{100h^{-1}\\,\\mbox{kpc}}\\right). \\end{equation} Lines of sight with angular separations of 3--10\\arcsec\\ are needed to probe the Jeans scale of the IGM.  While this scale is just larger than the cutoff for the typical Einstein radius of galaxy lenses \\citep*{schneider06}, new searches for binary quasars are revealing samples of a few to dozens with $\\Delta\\theta \\lesssim 10\\arcsec$ \\citep{hennawi06,hennawi09}.  This paper is organized as follows.  In \\S\\,\\ref{sec:sims}, we describe the SPH simulations used, as well as the artificial temperature-density relations we impose on the gas to isolate the effects of pressure support and thermal broadening.  In \\S\\,\\ref{sec:results} we discuss how the temperature-density relation affects the transverse coherence of the \\lya\\ forest, with particular focus on flux cross-correlation functions and conditional flux probability distributions.  We find that, as expected, the transverse coherence of the \\lya\\ forest across closely paired sightlines is dominated by the amount of pressure support in the absorbing gas.  These conclusions are summarized in \\S\\,\\ref{sec:conc}. In an Appendix and associated electronic tables, we provide \\lya\\ forest spectra extracted from our simulations at several transverse separations that can be used to create predictions tailored to specific observational analyses.  We note that Paper~I includes an extensive discussion of the physical structure of the \\lya\\ forest in these simulations, so in this paper we will focus only on those issues relevant to quasar pair observations.  All distances are given in comoving coordinates unless otherwise stated. ",
        "conclusions": "\\label{sec:conc} Recent efficient searches for binary quasars have yielded large samples of quasars with angular separations of $\\lesssim 10$\\arcsec\\ \\citep{hennawi06,hennawi09}.  The closely paired \\lyman\\ forest sightlines from such quasar pairs are ideal for studying the small-scale transverse structure of the intergalactic medium.  We have shown using a set of smoothed particle hydrodynamic simulations with different equations of state that the coherence of transverse structure at these scales is determined primarily by the level of pressure support, i.e.\\ the Jeans length, of the absorbing gas and is relatively insensitive to the amount of thermal broadening along the line of sight.  Given the surprisingly high temperatures implied by single-sightline analyses (see discussions in \\S\\ref{sec:intro} and Paper~I), it would be valuable to investigate the thermal state of the IGM by this largely independent method.  Flux correlation functions (measured by, e.g., \\citealt{rollinde03,coppolani06,dodorico06,marble08a}) and the two-point distributions of flux decrements or pixel ranks can all be used to distinguish among gas distributions with different temperatures like the fiducial and H4 models considered here.  Because the redshift ranges and pair separations will be dictated by the specifics of the available data sample, we provide in the form of electronic tables our simulated spectra along paired lines of sight from the two simulations, which can be used to generate predictions tailored to a particular data sample. These samples of spectra are described in \\S\\ref{sec:sims} and the caption to Table~\\ref{tab:rhot}.  We caution that the finite size of our simulation volume could cause statistical fluctuations and systematic effects on our predictions, especially at large $\\Delta r$ and $\\Delta v$; we will investigate this point in future work with larger simulations.  The large increase in known quasar pairs, the ability of large telescopes to measure \\lyman\\ absorption spectra of relatively faint background sources, which have a high surface density on the sky, and the massive quasar sample expected from the Baryon Oscillation Spectroscopic Survey \\citep{schlegel10} will change the \\lya\\ forest from a one-dimensional phenomenon to a three-dimensional phenomenon, opening new opportunities to constrain cosmology and the physics of the diffuse intergalactic medium."
    },
    "0910/0910.5828_arXiv.txt": {
        "abstract": "We report the discovery of a type IIn supernova in the low-metallicity dwarf galaxy J1320+2155, with an oxygen abundance 12+logO/H = 8.0$\\pm$0.2. This finding is based on SDSS (February 2008) and 3.5m Apache Point Observatory (February 2009) spectra taken one year apart, and on the observations that: the H$\\beta$ and H$\\alpha$ emission lines show broad components corresponding to gas expansion velocities of $\\sim$ 1600 km s$^{-1}$; the Balmer decrement is exceeedingly high: the H$\\alpha$/H$\\beta$ flux ratio, being more than 30, implies a very dense environment ($>$10$^7$ cm$^{-3}$); and the H$\\alpha$ broad luminosity decreases slowly, by only a factor of $\\sim$ 1.8 over the course of a year, typical of the slow luminosity evolution of a type IIn supernova. Several weak coronal lines of [Fe VII] and [Fe X] are also seen in the SDSS spectrum, implying ionization of the pre-shock circumstellar medium by shock-induced X-ray emission. The galaxy J1320+2155 is the first dwarf system ever to be discovered with a type IIn supernova exhibiting coronal lines in its spectrum. ",
        "introduction": "\\subsection{Variability of the broad  H$\\alpha$ luminosity}  The most notable features in the spectra of J1320+2155 are the strong and relatively broad H$\\alpha$ and H$\\beta$ emission lines. The parameters of these lines are shown in Table \\ref{tab3} for both the SDSS and high-resolution 3.5m APO spectra. Examination of  Table \\ref{tab3} shows a striking fact: the H$\\alpha$ to H$\\beta$ flux ratio is exceedingly high, being greater than 30, more than one order of magnitude greater than its recombination value. This implies that the broad hydrogen  emission comes from a high-density gas with an electron density in excess of 10$^7$ cm$^{-3}$. The large intensity of H$\\alpha$ relative to H$\\beta$ is due to electron collisional excitation in a very dense environment. The H$\\alpha$ luminosity derived from the SDSS spectrum has the high value $L$(H$\\beta$) = 6.43$\\times$10$^{40}$ erg s$^{-1}$. As discussed by \\citet{IT08}, such a high luminosity can be explained by the presence in the dwarf galaxy of either an AGN or a type IIp or type IIn supernova. Thanks to the APO spectra taken about one year after the SDSS spectrum, we can ascertain that the broad H$\\alpha$ luminosity has decreased by a factor of $\\sim$ 1.8 during that period. This drop in luminosity implies that the  H$\\alpha$ emission is produced by an expanding supernova envelope. We would expect a roughly constant luminosity in the case of an AGN and a more rapid fading in the case of a type IIp supernova.  \\subsection{A type IIn supernova}  We have no photometric observations of the supernova. We have therefore estimated its brightness from the continuum level of the SDSS spectrum obtained on February 11, 2008. In the wavelength range of $\\sim$ 4300 -- 4500\\AA\\ which corresponds to the $g$ filter, we obtain $g$ $\\sim$ 17.6 or about 0.3 mag brighter than the $g$ magnitude of J1320+2155 before the SN explosion, as derived from the SDSS $g$ image obtained on 12 March 2005 (Table \\ref{tab1}). Thus, the absolute $g$ magnitude of the supernova on the date of the SDSS spectral observations was $\\sim$ $-$15.6, as compared to the absolute $g$ magnitude of $\\sim$ $-$16.9 of the galaxy before the supernova event. The broad H$\\alpha$ and continuum flux distributions along the slit, peak at the same location in the 3.5m APO high-resolution spectrum, in the central part of the galaxy. This implies that the supernova is located in the nucleus of the dwarf galaxy. Note that the continuum level near H$\\alpha$ did not change over a period of one year, the time interval between the SDSS and APO observations. The constancy of the continuum level, the slow decline of the H$\\alpha$ luminosity and a nearly constant expansion velocity of $\\sim$ 1600 km s$^{-1}$ as measured from the FWHM of the broad component of the H$\\alpha$ emission line, are all characteristics of a  type IIn SN [SN 1988Z, \\citet{SS91}, \\citet{T93}, \\citet{S02}; SN 1995N, \\citet{F02}; SN 1997eg, \\citet{H08}; SN 2005ip, \\citet{S09}; SN 2007rt, \\citet{T09}].   \\subsection{Other lines with broad components}  Besides the H$\\alpha$ and H$\\beta$ lines, the He {\\sc i} $\\lambda$5876, $\\lambda$6678 and $\\lambda$7065 emission lines in the SDSS spectrum also show narrow and broad components. We use Gaussian profiles to fit both components. The narrow components fluxes of these lines are given in Table \\ref{tab2}. They show a striking fact: depending on the line, they are larger by a factor 1.5 -- 5 than the fluxes expected in a low-density ionized gas. This suggests a relatively dense environment in which the narrow He {\\sc i} lines are considerably enhanced by collisions in a dense pre-shock region of the supernova.  The fluxes relative to the broad H$\\beta$ flux and widths of the He {\\sc i} broad components are given in Table \\ref{tab4}. The broad component of the He {\\sc i} lines is also collisionally enhanced, but to a larger extent as compared to the narrow component, implying that they originate from larger density regions. In addition, a broad O {\\sc i} $\\lambda$8445 emission line is also detected in the SDSS spectrum. Its flux and width are also shown in Table \\ref{tab4}. All these lines have fluxes that are significantly smaller than the broad H$\\beta$ line flux. The errors of their fluxes and widths are consequently larger. Examination of Table \\ref{tab4} shows that the collisionally enhanced He {\\sc i} lines have widths that are lower but comparable within the errors to the widths of the broad hydrogen lines. These broad lines originate presumably in the post-shock circumstellar medium of the supernova.    \\subsection {Ionized gas diagnostics from nebular and auroral lines \\label{S3}}  Narrow nebular and auroral emission lines are frequently used for ionized gas diagnostics and for the determination of the electron temperature $T_e$, the electron number density $N_e$ and the element abundances in H {\\sc ii} regions of dwarf emission-line galaxies. Applying the procedures of e.g. \\citet{I06} and the semi-empirical method described by \\citet{IT07}, based on the fluxes of the strong nebular [O {\\sc ii}] $\\lambda$3727, [O {\\sc iii}] $\\lambda$4959, 5007 emission lines, we derive an oxygen abundance 12+logO/H = 7.68 for the ionized gas in J1320+2155. However, this technique is applicable only in the case of a low density gas ionized by stellar radiation. Usually, the electron density in the H {\\sc ii} region is derived from the nebular [S {\\sc ii}] $\\lambda$6717/$\\lambda$6731 line ratio. For J1320+2155, this ratio is high, $>$ 1, suggesting a low electron number density of the order of a few hundred cm$^{-1}$ \\citep{A84}. However, as discussed above, there are indications that the narrow line-emitting region in J1320+2155 has a high density (for example, the narrow He {\\sc i} emission lines are significantly enhanced by collisions), so that the low-density assumption for abundance determinations is not valid.  Additionally, the [O {\\sc iii}] $\\lambda$4363 and [N {\\sc ii}] $\\lambda$5755 auroral lines are relatively strong as compared to the nebular [O {\\sc iii}] $\\lambda$5007 and [N {\\sc ii}] $\\lambda$6583 emission lines. These lines are used for the determination of the electron temperature in H {\\sc ii} regions. Assuming a low density (100 cm$^{-3}$), we derive an electron temperature of $\\sim$32500 K from the [O {\\sc iii}] $\\lambda$4363/($\\lambda$4959+$\\lambda$5007) flux ratio. As for the [N {\\sc ii}] $\\lambda$5755/$\\lambda$6583 flux ratio, its value is so high that it cannot be explained by any temperature in the low density regime. Evidently, the regions that are emitting [O {\\sc iii}] and [N {\\sc ii}] are so dense that collisional deexcitation becomes important. This leads to an overestimate of the electron temperature and an underestimate of the abundances if a low density is assumed. Apparantly, these lines are emitted in the supernova pre-shock dense circumstellar region.  Since the critical densities for collisional deexcitation are significantly higher for auroral lines than for nebular lines, their emission comes preferentially from denser regions as compared to the emission of nebular lines. This fact may explain the apparent disagreement with the low density obtained from the low-ionization nebular [S {\\sc ii}] lines for which critical densities are low ($\\sim$ 10$^4$ cm$^{-3}$): these originate mainly in low-density regions.  We can estimate the electron number density by using the auroral-to-nebular flux ratios for the [O {\\sc iii}] and [N {\\sc ii}] emission lines and the equations from \\citet{A84} for temperature determination. These equations take into account collisional deexcitation. We find that for a reasonable range of temperatures, those varying between 10000 K and 15000 K, the electron number density must be in the range (1--2)$\\times$10$^{6}$ cm$^{-3}$ to account for the observed [O {\\sc iii}] $\\lambda$4363/($\\lambda$4959+$\\lambda$5007) and [N {\\sc ii}] $\\lambda$5755/$\\lambda$6583 flux ratios.  In such dense regions, element abundances cannot be derived because of the unknown spatial distribution of the gas density. We can obtain an estimate of the oxygen abundance of J1320+2155 by using the statistical luminosity-metallicity relation for dwarf irregular galaxies, e.g. the ones by \\citet{RM95} and \\citet{S89}. We find 12+logO/H = 8.0$\\pm$0.2, higher (by 0.3 dex) than the value obtained by the semi-empirical method, as expected because of the high density.   \\subsection{Coronal lines and diagnostics of high-ionization regions \\label{S4}}  The SDSS spectrum of J1320+2155 exhibits several narrow high-ionization coronal lines of [Fe {\\sc vii}] and [Fe {\\sc x}] (Fig. \\ref{fig2}). These lines have been seen before in spectra of type IIn SNe [SN 1988Z, \\citet{T93}, \\citet{S02}; SN 1995N, \\citet{F02}; SN 1997eg, \\citet{H08}; SN 2005ip, \\citet{S09}; SN 2007rt, \\citet{T09}] in spiral galaxies. However, their occurence is very rare. J1320+2155 is the first compact dwarf galaxy discovered with a type IIn SN spectrum showing coronal lines. If collisional ionization is the dominant mechanism, then the presence of these lines implies substantial amounts of gas at temperatures from several 10$^5$K to $\\sim$ 10$^6$K \\citep{B09} and pre-shock ionization of the circumstellar medium by X-ray emission. The gas can be cooler if photoionization is the main mechanism.  The critical densities for coronal lines are $\\sim$ 10$^{8}$ cm$^{-3}$ \\citep{NS82}, considerably higher than those for nebular lines of low-ionization species. The coronal lines can act thus as diagnostics of much hotter and denser gas in the supernova circumstellar medium. In particular, the observed [Fe {\\sc vii}] $\\lambda$3759/$\\lambda$6087 flux ratio of $\\sim$ 1.5 yields densities $<$ 10$^{7}$ cm$^{-3}$ for temperatures above $\\sim$ 100 000 K \\citep{NS82}. The high [Fe {\\sc vii}] $\\lambda$5159/$\\lambda$6087 flux ratio also implies that the electron density in the hottest regions is in the range $<$ 10$^{7}$ cm$^{-3}$. This range of number densities is broadly consistent with our estimates from the auroral-to-nebular flux ratios for the lower ionization species. It is also in the range of number densities obtained by \\citet{S09} for SN 2005ip, the SN with the spectrum richest in coronal lines known to-date, and taken several hundred days after the explosion. ",
        "conclusions": "}  We present SDSS (11 February 2008) and 3.5 m APO (23 February 2009) spectra of a newly discovered type IIn supernova in the compact dwarf galaxy J1320+2155.  We have obtained the following results:  1) The supernova was not present in the SDSS image obtained in March 2005. The host dwarf galaxy J1320+2155 on that image exibits a compact regular morphology with a size of $\\sim$ 1.5 kpc and an absolute SDSS $g$ magnitude of $-$16.89, for a distance of 92 Mpc. Its colors are typical of dwarf irregular galaxies. The oxygen abundance of the galaxy, estimated from the luminosity-metallicity relation for dwarf irregular galaxies, is 12+logO/H=8.0$\\pm$0.2.  2. From the light distribution in the slit, we infer that the supernova, which appeared prior to 11 February 2008, is in the central part of the galaxy. The broad widths of the H$\\alpha$ and H$\\beta$ emission lines, corresponding to velocities of $\\sim$ 1600 km s$^{-1}$ FWHM, an exceedingly high Balmer decrement (a  H$\\alpha$/H$\\beta$ flux ratio $>$ 30), the slow decline of the H$\\alpha$ broad luminosity over a time scale of $\\sim$ 1 year, all point to a type IIn SN. The spectrum also shows Helium lines that are broadened with velocities of $\\sim$ 1000 km s$^{-1}$ and that are collisionally enhanced.  3. Several narrow auroral and nebular emission lines of different low-ionization species are present in the SDSS spectrum. The auroral-to-nebular line flux ratios suggest that they originate in a dense ($\\sim$ 10$^{6}$ cm$^{-3}$) pre-shock circumstellar region.  4. Several narrow coronal emission lines of [Fe VII] and [Fe X] are detected in the SDSS spectrum, implying ionization of the pre-shock circumstellar medium by shock-induced X-ray emission.  There are reasons to suggest that type IIn SNe may occur in low-metallicity dwarf galaxies. The low metallicity lead to less cooling of the gas and hence to the formation of the more massive stars that are progenitors of this type of SNe. We plan more observations of the supernova in J1320+2155 to further constrain its properties."
    },
    "0910/0910.2639_arXiv.txt": {
        "abstract": "{ While the excess in cosmic-ray electrons and positrons reported by PAMELA and Fermi may be explained by dark matter decaying primarily into charged leptons, this does not necessarily mean that dark matter should not have any hadronic decay modes.  In order to quantify the allowed hadronic activities, we derive constraints on the decay rates of dark matter into $WW$, $ZZ$, $hh$, $q\\bar{q}$ and $gg$ using the Fermi and HESS gamma-ray data. We also derive gamma-ray constraints on the leptonic $e^+e^-$, $\\mu^+\\mu^-$ and $\\tau^+\\tau^-$ final states.  We find that dark matter must decay primarily into $\\mu^+\\mu^-$ or $\\tau^+\\tau^-$ in order to simultaneously explain the reported excess and meet all gamma-ray constraints.  } \\end{center} \\end{titlepage}  \\setcounter{page}{2} ",
        "introduction": "\\label{sec:1} The existence of dark matter (DM) has been firmly established by numerous observations, though the nature of DM mostly remains unknown. In particular, it is not known whether DM is absolutely stable or not.  If DM is unstable, it will eventually decay into lighter particles which may be observed as an excess in the cosmic-ray spectrum.  If DM is related to new physics which appears at the weak scale, it is natural to expect that the DM mass is in the range ${\\cal O}(100)\\,{\\rm GeV}$--${\\cal O}(10)$\\,TeV.  However, the longevity of DM whose mass is of the weak scale is a puzzle and calls for some explanation.  The (quasi)stability may be the result of a discrete symmetry or extremely weak interactions. For instance, in a supersymmetric (SUSY) theory, the lightest SUSY particle (LSP) is stable and therefore a candidate for DM if R-parity is an exact symmetry. However, R-parity violation may be a common phenomenon in the string landscape~\\cite{Kuriyama:2008pv}, in which case the LSP DM is unstable and eventually decays into Standard Model (SM) particles. On the other hand, if DM is in a hidden sector which has extremely suppressed interactions with the SM sector, the only way to probe DM may be to look in the cosmic rays for signatures of its decay products.   Recently, much attention has been given to the electron/positron excess reported by PAMELA~\\cite{Adriani:2008zr}, ATIC~\\cite{ATIC}, PPB-BETS~\\cite{Torii:2008xu} and Fermi~\\cite{Abdo:2009zk}.  The excess clearly suggests that we need to modify our current understanding of the production/acceleration/propagation of cosmic-ray electrons and positrons. Of the many explanations proposed so far for this excess, DM decay or annihilation remains an exciting possibility.  In order to account for the excess by DM decay/annihilation, DM should mainly produce leptons with suppressed hadronic branching ratios, otherwise the anti-protons produced would likely exceed the observed flux~\\cite{Adriani:2008zq}. Further model-independent analysis also revealed that, if one requires that the PAMELA/Fermi excess be explained by DM annihilation, the gamma-rays accompanying the lepton production exceeds the observed flux unless a significantly less steep DM profile is assumed~\\cite{Bertone:2008xr,meade:fermi}.  Thus, the leptophilic decaying DM scenario has recently gained momentum~\\cite{Chen:2008yi,Chen:2008dh,Yin:2008bs, Ishiwata:2008cv,Shirai:2009kh,Ibarra:2009bm,Ibe:2009en}.  Leptophilic decaying DM models can be broadly divided into two categories.  One is such that DM first decays into additional light particles, which subsequently decay into muons or electrons, while the decays into hadrons are forbidden by kinematics~\\cite{Cholis:2008vb}. The other is such that the DM particle couples mainly to leptons due to symmetry~\\cite{Chen:2008dh} or geometric setup~\\cite{Okada:2009bz}. While the hadronic activities are absent in the former case, it is model-dependent to what extent the DM is leptophilic in the latter case.  One example is the hidden gauge boson decaying into the SM particles through a mixing with a U(1)$_{B-L}$ gauge boson~\\cite{Chen:2008yi}; the DM is certainly leptophilic in the sense that it mainly decays into leptons, but a certain amount of quarks are also produced.  The purpose of this paper is to study the current constraints on the hadronic and leptonic decay of DM. The anti-proton flux is known to provide a tight constraint on the hadronic activities, but there are large uncertainties in the propagation~\\cite{Hisano:2005ec}. The other constraint comes from gamma-ray observation. In contrast to charged cosmic-ray particles, gamma-rays travel undeflected and there is no uncertainty in the propagation; the main uncertainty is the dark matter profile.  Furthermore, the Fermi satellite has been measuring gamma-rays with both unprecedented precision and statistics, and we can expect a significant improvement over EGRET data~\\cite{Sreekumar:1997un,Strong:2004ry}.  In this paper we will derive constraints on the partial decay rates of DM into $WW$, $ZZ$, $hh$, $q\\bar{q}$ and $gg$ as well as $e^+e^-$, $\\mu^+\\mu^-$ and $\\tau^+\\tau^-$ using the Fermi~\\cite{Fermi:ML, Fermi:iso} and HESS~\\cite{HESS:GC} data.  The bounds obtained in this paper are not only generic but also can be used to know what branching ratios are allowed in decaying DM models which account for the PAMELA/Fermi cosmic-ray electron/positron excess. ",
        "conclusions": "\\label{sec:5} In this paper we have derived constraints on the partial decay rates of DM into $WW$, $ZZ$, $hh$ and $q\\bar{q}$ as well as $e^+e^-$, $\\mu^+\\mu^-$ and $\\tau^+\\tau^-$ using the gamma-ray data observed by the Fermi LAT and HESS. One of the merits of using gamma-ray is that the predicted flux does not depend on the diffusion model in the Galaxy, in contrast to charged cosmic-rays. The constraints derived in this paper provide implications for DM model-building attempting to account for the PAMELA/Fermi excess.  According to our results, the allowed hadronic branching ratio is of ${\\cal O}(10)$ \\%.  We have applied the result to a DM model based on the hidden gauge boson decaying through a mixing with the U(1)$_{B-L}$, and found that the model is now marginally excluded by the Fermi gamma-ray observation. The allowed hadronic branching ratio is about $30\\%$ at the reference point shown as a star in Fig.~\\ref{fig:bminusl}.  Gamma-ray observations have the power to constrain the properties of DM.  Thanks especially to the Fermi LAT data, which provides greater accuracy and statistics over the experiments in the past, the gamma-ray constraint has become as tight as or even tighter than current neutrino and antiproton constraints.  Further observation by the Fermi satellite will hopefully shed light on the dark sector of our Universe."
    },
    "0910/0910.1310_arXiv.txt": {
        "abstract": "We use the number density of peaks in the smoothed cosmological density field taken from the 2dF Galaxy Redshift Survey to constrain parameters related to the power spectrum of mass fluctuations, $ n $ (the spectral index), d$n$/dln$k$ (rolling in the spectral index), and the neutrino mass, $m_{\\nu}$. In a companion paper we use N-body simulations to study how the peak density responds to changes in the power spectrum, the presence of redshift distortions and the relationship between galaxies and dark matter halos. In the present paper we make measurements of the peak density  from 2dF Galaxy Redshift Survey data, for a range of smoothing filter scales from $4-33 \\hmpc$. We use these measurements to constrain the cosmological parameters, finding $n=1.36^{+0.75} _{-0.64}$, $m_{\\nu} < 1.76 $eV, $d$n$/dln$k$=-0.012^{+0.192}_{-0.208}$, at the 68 \\% confidence level, where $m_{\\nu}$ is the total mass of three massive neutrinos. At 95\\% confidence  we find $m_{\\nu}< 2.48$ eV. These measurements represent an alternative way to constrain cosmological parameters to the usual direct fits to the galaxy power spectrum, and are expected to  be relatively insensitive to non-linear clustering evolution and galaxy biasing. ",
        "introduction": "The space density of peaks in a cosmological density field smoothed with a filter is sensitive to the shape of the linear power spectrum of mass fluctuations, even for filter sizes where the fluctuations are in the non-linear regime (Croft \\& Gazta\\~naga 1998). In De \\& Croft (2007), hereafter Paper I we explored the sensitivity of the peak density to parameters related to the initial power spectrum, as well as redshift distortions and variations in the galaxy halo occupation distribution (e.g., Berlind \\& Weinberg 2002). The theory of peaks in a Gaussian density field was set out in detail by Bardeen \\etal (1986, hereafter BBKS). The relationship between the peak density and power spectrum in BBKS can be used to explore  such parameters as the spectral index and its dependence on scale (Kosowsky and Turner, 1991) along with the neutrino mass. In this paper we use data from the 2dF Galaxy Redshift Survey (hereafter 2dFGRS, Colless \\etal 2001), to constrain cosmological parameters using the peak density.  In a recent analysis of the 2dFGRS final data set Cole \\etal (2005)  employed a direct Fourier method to compute the power spectrum. These authors put constraints on several parameters using the directly measured power spectrum shape. They assumed a primordial form for P(k) with $n=1$ and $h=0.72$ along with a negligible neutrino mass. This gave preferred values of $\\Omega _{m}h=0.168 \\pm 0.016$ and a baryon fraction of $\\frac{\\Omega_{b}}{\\Omega_{m}}=0 .185 \\pm 0.046$( 1$\\sigma$ errors). This analysis therefore implies a significantly lower mass density compared to the $\\Omega_{m}=0.3$ often taken as standard. In combination with CMB data from WMAP (Spergel \\etal 2003) Cole \\etal find $\\Omega_{m}=0.231 \\pm 0.021$ Also on large scales some  evidence was seen by Cole \\etal  of the baryon oscillations predicted by CDM models. Our present work is complimentary to this large scale linear theory analysis, with the peak density enabling constraints to be placed on the linear power spectrum and cosmological parameters from data on smaller scales.   This paper is the second in a series. In paper I, we  examined the number density of peaks in cosmological simulations and compared to results from linear peak theory (Bardeen \\etal 1986, hereafter BBKS) over a range of density field smoothing filter scales. This provided knowledge of the length scales for which agreement between theory and simulation can be expected. The dark matter simulations used to compare to peak theory showed good agreement for filter scales between 3-30$\\hmpc$. This was assuming that the mean interparticle separation was smaller than the filter scale (for agreement at better than the $5 \\%$ level). In addition to simple tests using simulations, we created galaxy catalogues using lists of dark matter halos and the Halo Occupation distribution formalism (Zheng \\etal 2005.) Good correspondance was found between the peak density measured from  the galaxies and dark matter for filter scales $>4 \\hmpc$.  The peak density can be used to constrain cosmology through its dependence on the power spectrum of mass fluctuations, $P(k)$. In particular we work with the asymptotic number density of maxima, i.e. the  number density of peaks of all heights and which is found by BBKS to be: \\begin{equation} \\langle n_{\\rm pk} \\rangle = \\frac{29-6\\surd6}{5^{\\frac{3}{2}}2(2\\pi)^{2}R_{\\ast}^{3}}, \\label{npke} \\end{equation} where $R_{\\ast}$ is defined to be the following ratio of moments of the power spectrum: \\begin{equation} R_{\\ast}=\\sqrt{3}\\frac{\\sigma_{1}}{\\sigma_{2}}. \\end{equation} Here \\begin{equation} \\sigma_{j}^{2} =\\frac{\\int k^{2}P(k)k^{2j} dk}{2\\pi^{2}}. \\end{equation} The scale dependence of $P(k)$ is probed by  smoothing the density field with Gaussian filter with comoving radius $r_{\\rm f}$, i.e. multiplying $P(k)$ by $e^{-(kr_{\\rm f})^{2}}$. Croft \\& Gazta\\~naga (1998) found that there is an simple relationship that holds to high accuracy between the local slope of the power spectrum $n$ at wavenumber $k$ and the filter scale $r_{f}\\sim0.4 \\pi/k$. The peak number density is therefore sensitive to the power spectrum slope (but not its amplitude) and the parameters which govern it can be tested.   Paper I  focussed on several such aspects of the power spectrum such as the neutrino mass, the possible rolling of the spectral index $n$, and redshift distortions. It was seen that upon increasing the neutrino mass the peak density decreases, as it does when the rolling of the spectral index is made more negative. The effect of redshift distortions was also quantified, showing that they act to suppress the number density of peaks on small scales. This information is useful for the present paper which deals with the peak density in redshift space measured from an observed galaxy catalogue.     Our approach in this paper is to first test our recovery of the true peak density using mock catalogues derived from simulations that have had the 2dF masks and selection function applied to them. We then apply the same estimation techniques to the 2dF data, using the recovery from mock data to estimate the reliability of the method. The 2dF peak density as a function of filter scale is then used to constrain cosmological parameters.  The plan for the paper is as follows: In  Section 2 we  describe the simulated observations (mock catalogues), including the detailed reasoning behind their use. We give an overview of the 2dFGRS and in Appendix A describe how the mock catalogues were generated. The latter includes a description of the selection function, magnitude limits and geometrical limitations of the survey and how they were applied to N-body simulations to produce mock catalogues. In Section 3 we detail tests of the process of finding peaks in the simulations and mock catalogues, including selection of the volume limits for a given filter size and how we deal with incomplete sky coverage. In Section 4 we  describe the calculation of the observed peak density from the 2dF survey data using the completeness functions calibrated from mock catalogues. We also describe various sources of uncertainty, and how they propagate into our evaluation of cosmological parameters. In Section 5 we  present the best fit values of $m_{\\nu}$, $n$ and d$n$/dln$k$ including error estimates, using Markov Chain Monte Carlo analysis. In Section 5 we conclude and discuss possible future work. ",
        "conclusions": "By measuring the space density of peaks in the smoothed 2dF galaxy survey density field as a function of filter scale, we have constrained the values of the cosmological power spectrum parameters d$n$/dln$k$, $n$ and $\\Omega_{\\nu}$. We find values that are consistent within 1 $\\sigma$ of those measured from  other cosmological observables (e.g., Dunkley \\etal 2009, Seljak \\etal 2005), and are in support of the standard $\\Lambda$CDM scenario. In linear theory, the peak density is not affected by the amplitude of fluctuations. As simulation tests have shown that this is also true in the non-linear regime, for filter scales as small as $3 \\hmpc$, this means that our constraints on cosmological parameters come entirely from the shape of the power spectrum. Because of degeneracies in the shape of the power specrum with cosmological parameters it was necessary to assume prior values for some parameters ($\\Omega_{m}, \\Omega_{b}$ and $H_{0}$) from the WMAP (Spergel \\etal 2003) results in order to derive our own limits.   Various sources of systematic errors are possible, when measuring peak statisics, for example those related to the fact that the peak density in a smoothed field is sensitive to boundary effects, which are difficult to correct for (unlike the case of the correlation function, for example.) In the present work, we have made conservative cuts, ignoring space close to the survey boundaries and validated these using mock catalogs to test both the overall recovery of the peak density and the recovery of individual peaks. An additional source of error comes from the fact that redshift distortions tend to merge peaks together which are separate in real space, as was seen in paper I. We have dealt with this effect, by including a power spectrum suppression term due to the small scale random velocity dispersion (see paper I for details and tests). This is the one part of the analysis for which peak finding may be most sensitive to the relationship between mass and light. Our tests using the halo occupation distribution in Paper I have show that this is unlikely to be a problem at the level of current observational uncertainties.  Our constraints have come from an analysis of  filter lengthscales between $4-30 \\hmpc$. This approach is therefore complementary to much work in  large scale structure which has been done by analyzing the galaxy power spectrum on somewhat larger length scales (wavenumber $k< 0.1 \\invhmpc$, corresponding to wavelengths $\\sim \\pi/k > 30 \\hmpc$). On these larger scales, galaxy fluctuations are assumed to be linearly related to mass, and direct measurements of the galaxy power spectrum are used to constrain the mass power spectrum and hence cosmological parameters (e.g., Tegmark \\etal 2004).  Recently, it was shown by Sanchez \\& Cole (2007) that shapes  of the power spectra measured from the 2dF and SDSS galaxy surveys (e.g., Cole \\etal 2005, Tegmark \\etal 2004) are not in agreement, due to the $r$-band selected SDSS galaxies having a stronger scale-dependent bias. This results in differences in the cosmological parameters inferred from the two surveys which are larger than the quoted measurement uncertainties. From the tests in CG97 and Paper I, we expect that our peak based measurement to be more robust  both to differences in galaxy bias and non-linear evolution. In the future it will be useful to compute the peak number density from the SDSS survey data, in order to check consistency. The larger size of the SDSS final dataset will enable the use of larger filter scales and reduce the error bars on parameters."
    },
    "0910/0910.1060_arXiv.txt": {
        "abstract": "We present {\\it Chandra} X-ray observations of the powerful radio galaxy 3C\\,171, which reveal an extended region of X-ray emission spatially associated with the well-known 10-kpc scale optical emission-line region around the radio jets. We argue that the X-ray emission comes from collisionally ionized material, originally cold gas that has been shock-heated by the passage of the radio jet, rather than being photoionized by nuclear radiation. This hot plasma is also responsible for the depolarization at low frequencies of the radio emission from the jet and hotspots, which allows us to estimate the magnetic field strength in the external medium. We show that it is likely that both the cold emission-line gas and the hot plasma in which it is embedded are being driven out of the host galaxy of 3C\\,171 at supersonic speeds. A significant fraction of the total energy budget of the central AGN must have been expended in driving this massive outflow. We argue that 3C\\,171, with its unusual radio morphology and the strong relation between the jet and large amounts of outflowing material, is a member of a class of radio galaxies in which there is strong interaction between the radio jets and cold material in the host galaxy; such objects may have been very much more common in the early universe. ",
        "introduction": "In powerful radio galaxies, extended optical emission-line regions (EELR) are often found to be aligned with the axis defined by the extended radio emission (e.g., \\citealt{mvbs87}; \\citealt{m93}; \\citealt*{mbs96}). A long-standing question \\citep[e.g.,][]{bh89,tmrd98} is whether these EELR are ionized by photons from the nucleus or by shocks driven by the jets. Shocks are part of the standard model for these powerful objects \\citep[e.g.,][]{bc89} but direct observational evidence for them has been somewhat elusive, and to date has mostly come from the X-ray band (e.g., \\citealt{kvfj03}; \\citealt{omwk06}; \\citealt*{wsy06}). However, in the most powerful radio sources, normally found at high redshift, there is often evidence from one or more of the line kinematics, the line ratios or the relationship with the radio morphology that the optical emission lines are either ionized or at least strongly affected by jet-driven shocks, which requires a direct interaction between the jets and the {\\it cold} ($T \\la 10^4$ K) phase of the IGM of the host galaxy (e.g., \\citealt*{stb02}; \\citealt{nldg08}). Photoionization presumably still operates, as it must in any source with a radiatively efficient accretion flow, but the interaction between jets and cold material dominates the energetics of the EELR. Understanding such interactions is important in models in which accretion onto the central black hole affects the properties of galaxies hosting powerful AGN (`feedback'), while physical conditions in and kinematics of the EELR in these systems may give us important information about the energetics of the radio source as a whole. Unfortunately, in nearby, comparatively low-luminosity radio galaxies, which have the advantage that they can be studied with high spatial resolution, EELR are generally small, centrally peaked and can adequately be described by central-illumination models; they are plausibly just powerful versions of the ionization cones seen in local radio-quiet AGN such as Seyfert galaxies \\citep[e.g.,][]{bh89}. Thus the study of EELR in nearby radio-loud sources does not in general shed any light on jet-gas interactions.  There is evidence, however, that a small minority of low-redshift radio galaxies are also more adequately described by a shock-ionization model. One good candidate for a shock-ionization source at low redshift is 3C\\,171, an unusual $z=0.2384$ radio galaxy. Low-resolution radio imaging \\citep*{hvm84,b96} shows that in its inner regions it is similar to a normal FRII, but low-surface-brightness plumes instead of normal lobes extend north and south from the hotspots. High-resolution images show knotty jets connecting the radio core with the hotspots \\citep{hapr97}. 3C\\,171 is associated with a prominent EELR, elongated along the jet axis, and the emission-line gas is brightest near the hotspots \\citep{hvm84}, suggesting that the radio-emitting plasma is responsible for powering the emission-line regions. \\cite{catr98} made William Herschel Telescope observations of the emission-line regions and showed that their ionization states are most consistent with a shock-ionization model, with the disturbed kinematics and high-velocity line splitting being indicative of direct driving of motions in the line-emitting gas by the jets. The velocity centroids of the lines observed are blueshifted at the level of a few hundred km s$^{-1}$ relative to the systemic velocity in the east and similarly redshifted in the west, consistent with a jet-driven outflow (note that the two-sided appearance of the radio jets implies, given the normal assumption of relativistic jet speeds, that the jets are not far from the plane of the sky, in which case these radial velocities may substantially underestimate the outflow speed of the cold gas). The results of \\citeauthor{catr98} were confirmed by later studies with integral-field spectroscopy and with the {\\it Hubble Space Telescope} ({\\it HST}) \\citep{st03,totw05}. In terms of its apparent strong interactions between jet and cold IGM, 3C\\,171 may be a low-redshift analog of EELR around radio galaxies at high redshift, where cold gas is expected to be more abundant in radio source host galaxies.  3C\\,171 is also the only radio galaxy to date in which physical conditions in the EELR have been probed with detailed radio depolarization studies. \\citet{hvm84} pointed out the depolarization of the radio source near the hotspots, while \\citet{hapr97} showed that the depolarization was spatially very closely connected with the EELR. More recently \\cite{h03} (hereafter H03) used high-frequency radio polarimetry to measure the depolarization and Faraday rotation in the source as a function of position. Radio depolarization of the type seen in 3C\\,171 comes about as a result of unresolved structure in the Faraday rotation due to magnetoionic material in front of the source. As Faraday rotation depends on the electron density and magnetic field strength integrated along the line of sight, measurements of depolarization provide a measurement of the product of electron density and magnetic field strength in the external medium \\citep{b66} which is essentially independent of the other physical conditions in that medium, such as temperature, so long as it is ionized. The key result of H03 was that the line-emitting material of the EELR, though spatially coincident with the observed depolarization, cannot itself be responsible for the depolarization (given the physical parameters deduced from the observations and the known physical parameters of line-emitting clouds, the clouds have far too low a covering factor). Instead the EELR-emitting clouds must be embedded in a more diffuse, and presumably hotter, medium whose parameters were constrained by the radio data.  X-ray studies of 3C\\,171 allow a search for this diffuse phase of the external medium of the source. At the time that H03 was writing the only imaging X-ray data available were a {\\it ROSAT} HRI observation that had previously been shown \\citep{hw99} to give no detection of extended emission. Based on this and on assumed values for the characteristic magnetic field strength in the external medium, H03 placed some limits on the properties of the hot phase which suggested that an X-ray detection with deeper observations was unlikely. However, more recently, a short {\\it Chandra} observation taken on 2007 Dec 22 (obsid 9304), with an effective sensitivity a factor of a few deeper than that of the {\\it ROSAT} observations and a much higher spatial resolution, (Massaro \\etal\\ in prep) showed extended X-rays that clearly correlated very well with the observed region of radio depolarization, implying that the depolarizing medium could indeed be imaged directly.  The results of the {\\it Chandra} snapshot motivated us to make the deep {\\it Chandra} observation that is the subject of the present paper. In this paper we first describe the deep X-ray observation (Section 2) and the principal imaging and spectroscopic results (Section 3). In Section 4 we first argue that the extended X-ray emission is definitely thermal in origin and associated with the depolarizing medium and then go on to discuss the implications of this model for the dynamics and energetics of the source. Our conclusions are summarized in Section 5.  Throughout the paper we assume a cosmology with with $H_0 = 70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\rm m} = 0.3$ and $\\Omega_\\Lambda = 0.7$. This gives a luminosity distance to 3C\\,171, at $z = 0.2384$, of 1194 Mpc and an angular scale of 3.77 kpc arcsec$^{-1}$. ",
        "conclusions": "Our main results on 3C\\,171 can be summarized as follows.  \\begin{enumerate} \\item We find an extended region of X-ray emission coincident with the well-studied extended optical emission-line region in 3C\\,171. \\item In what we feel is the most plausible interpretation, the X-ray emission is thermal and is due to collisionally ionized hot plasma in which the cold clouds of optical line-emitting material are embedded. On these scales photoionization models require very high continuum luminosities and are hard to render consistent with the dynamics of the source and the observed radio depolarization. \\item In the collisionally ionized interpretation, the X-ray emitting material is capable of producing the depolarization of the radio emission from the jets and hotspots studied by H03. The physical conditions we estimate from the X-ray observation give rise to very reasonable constraints on the magnetic field strength in the medium external to the radio source, $2\\ \\mathrm{nT} < B < 9\\ \\mathrm{nT}$, and a scale for the magnetic field reversals in the range $0.2\\ \\mathrm{kpc} < d < 1.0\\ \\mathrm{kpc}$. \\item If the hot plasma is moving with the cold line-emitting material, then it dominates the energetics of the outflow. The total energy budget required depends on the outflow speed adopted, but lies in a range between a few $\\times 10^{43}$ and a few $\\times 10^{44}$ erg s$^{-1}$, comparable to the bolometric radiative luminosity of the active nucleus. \\end{enumerate}  These results are significant for two reasons. Firstly, the very good agreement between the physical conditions implied by the optical line emission, the X-rays and the radio depolarization gives us some confidence that we have at least a basic picture of the complex multi-phase medium (consisting of at least warm gas, hot plasma and magnetic fields) in which the jet is embedded. Over the coming decades, our ability to study polarization structures in radio galaxies will be greatly improved, while our ability to image them with high resolution in X-rays may unfortunately be significantly diminished. It is therefore important to develop methods of inferring physical conditions from radio polarization (and optical emission-line) studies alone, and our work shows that this can be done. Secondly, they imply that many of the sources in which extended emission-line regions have been detected may also have energetically dominant outflows of hot plasma; if confirmed by X-ray studies of suitable targets, this makes optical emission-line studies central to an understanding of the interaction between radio galaxies and the cold gas in their environments."
    },
    "0910/0910.4968_arXiv.txt": {
        "abstract": "We present new observational results on the kinematical, morphological, and stellar population properties of a sample of 21 dEs located both in the Virgo cluster and in the field, which show that 52$\\%$ of the dEs i) are rotationally supported, ii) exhibit structural signs of typical rotating systems such as discs, bars or spiral arms, iii) are younger ($\\sim3$\\,Gyr) than non-rotating dEs, and iv) are preferentially located either in the outskirts of Virgo or in the field. This evidence is consistent with the idea that rotationally supported dwarfs are late type spirals or irregulars that recently entered the cluster and lost their gas through a ram pressure stripping event, quenching their star formation and becoming dEs through passive evolution. We also find that all, but one, galaxies without photometric hints for hosting discs are pressure supported and are all situated in the inner regions of the cluster. This suggests a different evolution from the rotationally supported systems. Three different scenarios for these non-rotating galaxies are discussed (in situ formation, harassment and ram pressure stripping). ",
        "introduction": "Dwarf galaxies ($M_B>-18$)  are the most numerous objects in the  nearby Universe. They have gained importance since hierarchical models proposed dwarfs as the building blocks of massive galaxies \\citep[e.g.,][]{WR78,WF91}. These systems can be divided in star forming (Im, BCD, Sd,...) and quiescent (dE, dS0) dwarfs, the latter being the dominant population in clusters, the former  the most common in the field \\citep{Sand85,FB94,Blant05,Crot05}. The study of the population of low-mass galaxies and of the mechanisms leading to the strong morphology segregation between the different types of dwarfs is fundamental to understand the assembly and evolution of the overall population of galaxies.   In the case of dEs, various scenarios have been suggested to explain their formation. Whether they are the low luminosity extension of the giant ellipticals (Es) or they are the result of different formation and evolution processes is still an open question \\citep{FB94}. It has been proposed that dEs were late-type systems where the interstellar medium (ISM)  was swept away by the kinetic pressure due to supernova explosions (\\citet{YosAri87}; see however \\citet{ST01}). Other theories suggested that late-type spirals stopped their star formation once their ISM was removed during their interaction with the environment and evolved into quiescent dwarfs. The existence of a morphological segregation effect on the dwarf galaxy population \\citep{FB94} favors this second scenario. Different processes could be at the origin of gas removal in clusters: interaction with the intergalactic medium (IGM) \\citep[e.g.,][]{VZ04}, as ram-pressure stripping, galaxy-galaxy interactions \\citep[e.g.,][]{ByrdValt90} and harassment \\citep[e.g.,][]{Moore98, Mast05}. These mechanisms are  able to reproduce some of the observed properties of local dEs in clusters. Structural parameters, such as surface brightness, and spectrophotometric properties, such as stellar populations, are easily reproduced after a ram pressure event \\citep{Boselli08a,Boselli08b}. The detailed study of \\citep[][]{Lisk06b,Lisk06a,Lisk07} shows that there exists a population of galaxies that have properties in between those of dEs and star forming dwarfs. In these objects, which are classified as dEs, some remains of spiral discs, such as spiral arms or irregular features, are still visible. If dEs are formed from star forming galaxies through ram pressure stripping, we expect that the angular momentum of the parent galaxies should be conserved, while in the case of multiple dynamical perturbations the system is rapidly heated and the rotation is lost. Measuring the kinematics of dEs is therefore an important test to understand their origin.   In the last decade considerable efforts have been made to measure the kinematic properties of dE in Virgo, the closest rich cluster \\citep[e.g.,][]{Ped02,Geha02,Geha03,VZ04}. These works have shown the existence of both rotating and pressure supported systems. This fact, together with the varieties of morphologies found by \\citet{Lisk06b} makes it clear that the formation of dEs might rather be complex. Would every different population of dEs have a separate formation process? To answer this question we started a kinematical survey of dE in the Virgo cluster.This is the first of a series of papers devoted to the study of the kinematics and stellar populations of dEs in the Virgo cluster. ",
        "conclusions": "It is believed that the difference between fast and slow rotators in Es is due to a different formation origin, with slow rotators, the most massive objects, principally located in dense environments and resulting from major mergers, while fast rotators, being less massive and more gas$-$rich \\citep{Em07}, have a star formation history that is more extended in time. For dwarf galaxies this segregation in mass is difficult to be appreciated because of their small dynamical range (8.8$<L_H/L_{H\\odot}<$9.7 in our sample). Moreover, models indicate that the probability that dEs have undergone a major merging event is almost null \\citep{deLucia06}. This is consistent with the fact that dEs are on average much younger than Es \\citep{Mich08}, whose formation happened probably at $z>$2 \\citep{Renz06}. If their star formation activity was quenched in recent epochs, this would hardly be due to major merging events since the high $\\sigma$ of clusters already formed prevented strong dynamical interactions \\citep{BG06}.  Dwarf elliptical systems might have a different origin according to their structural, spectrophotometric and kinematical properties. Although harassment would be able to explain the origin of the pressure supported dwarfs, it is not likely to explain all the properties of rotationally supported systems ($(v_{max}/\\sigma)^*>$0.8). Harassment has long time scales since it needs multiple encounters to be efficient in removing gas and stars and quenching the star formation \\citep{BG06}, a scenario inconsistent with the observations which indicate that the star formation was active down to recent epochs, as the luminosity-weighted ages derived by \\citet{Mich08} suggest (see Figure \\ref{f3}). Furthermore, dynamical interactions heat the system removing rotation and enhancing velocity dispersion. In the simulations of \\citet{Mast05} those galaxies that keep rotation after the harassment still retain an important spiral structure, yet they can not be classified as ellipticals. Moreover harassment is more efficient in the core of the clusters (in the inner 100kpc), while our result indicates that these rotationally supported systems are located mostly in the outskirts. All these evidences are however consistent with a recent ram pressure stripping event as proposed by \\citet{Boselli08a,Boselli08b}. Indeed, ram pressure would easily remove the total gas content of the infalling low luminosity, late-type galaxies, quenching their star formation, but conserving, at least on short time scales, their angular momentum and therefore their rotation. The spectrophotometric and structural properties of dE in Virgo are consistent with a recent ($\\le$2 Gyr) interaction with  the cluster hot IGM \\citep{Boselli08a,Boselli08b}. Consistent with observations \\citep{CK08}, hydrodynamical simulations indicate that ram pressure stripping can be efficient up to the virial radius, or even more outside in dwarf systems with shallow potential wells as those analysed here \\citep{Roed08}.    Pressure supported galaxies ($(v_{max}/\\sigma)^*<0.8$) might have a different origin:  1) they might be galaxies formed in situ through the isotropic collapse of the gas at early epochs, thus not transformed by the environment, and be the extension of Es, with the exception that they probably did not undergo major merging events \\citep{deLucia06}. Indeed they have old stellar populations, they do not present spiral structures \\citep{Lisk06b} and they are virialised within the cluster, thus are members since early epochs ~\\citep{Conselice01}. 2) They can be star forming systems that entered into the cluster in the early epochs, maybe through the accretion of groups where preprocessing was active, and later modified by galaxy harassment or 3) they are star forming systems that entered into the cluster several Gyr ago, where ram pressure, although less efficient than today, because of the lower density of the IGM and of the smaller velocity dispersion of the cluster still in formation, had the time through multiple cluster crossing to remove the gas and stop the star formation activity.  The lack of supply of new generations of stars on the plane of the disc increases $\\sigma$  on long time scales, while a decrease of the rotation would in any case require dynamical interactions. The 2 $\\sigma$-dominated disc dwarfs (VCC1183, VCC1910) could be galaxies in the last stages of their transformation into dEs.   The analysis done so far, although crucial for understanding the origin of the rotationally supported systems, is still insufficient to explain the origin of the $\\sigma$ supported galaxies. It is probable that harassment and ram pressure  were acting together with a relative weight that might have changed with time. Observations of high redshift clusters, mainly populated by massive, red sequence galaxies at late epochs \\citep{deLucia06} consistently suggest that harassment and ram pressure are the most probable effects. It is however clear that a complete study of the kinematical, structural and spectrophotometric properties of early and late type galaxies in clusters combined with model predictions, is a powerful way to study the role of the environment on their evolution, a work we plan to do in the near future.    In summary, we report the first clear evidence of a correlation between the presence of rotationally supported dEs and the distance to the center of the Virgo cluster. We have seen that the further away a dwarf galaxy is from M87, the larger its rotation is and the younger it appears. In the outer parts many of them show spiral features beside their elliptical appearance. These kinematical data are consistent with the idea of ram pressure stripping transforming dwarf star forming galaxies into dEs. Dwarf ellipticals that are pressure supported and show no sign of spiral features are likely to be late-type systems that entered in the cluster at early epochs and were later transformed by a combined effect of multiple encounters with nearby companions and the interaction with the hot and dense IGM. They might also be originally formed in clusters from an isotropic collapse of the primordial gas cloud."
    },
    "0910/0910.3065_arXiv.txt": {
        "abstract": "Three-dimensional numerical simulations show that large-scale latent heating resulting from condensation of water vapor can produce multiple zonal jets similar to those on the gas giants (Jupiter and Saturn) and ice giants (Uranus and Neptune).  For plausible water abundances (3--5 times solar on Jupiter/Saturn and 30 times solar on Uranus/Neptune), our simulations produce $\\sim20$ zonal jets for Jupiter and Saturn and 3 zonal jets on Uranus and Neptune, similar to the number of jets observed on these planets. Moreover, these Jupiter/Saturn cases produce equatorial {\\it superrotation} whereas the Uranus/Neptune cases produce equatorial {\\it subrotation}, consistent with the observed equatorial jet direction on these planets. Sensitivity tests show that water abundance, planetary rotation rate, and planetary radius are all controlling factors, with water playing the most important role; modest water abundances, large planetary radii, and fast rotation rates favor equatorial superrotation, whereas large water abundances favor equatorial subrotation regardless of the planetary radius and rotation rate.   Given the larger radii, faster rotation rates, and probable lower water abundances of Jupiter and Saturn relative to Uranus and Neptune, our simulations therefore provide a possible mechanism for the existence of equatorial superrotation on Jupiter and Saturn and the lack of superrotation on Uranus and Neptune. Nevertheless, Saturn poses a possible difficulty, as our simulations were unable to explain the unusually high speed ($\\sim400\\rm\\,m\\,s^{-1}$) of that planet's superrotating jet. The zonal jets in our simulations exhibit modest violations of the barotropic and Charney-Stern stability criteria. Overall, our simulations, while idealized, support the idea that latent heating plays an important role in generating the jets on the giant planets. ",
        "introduction": "The question of what causes the prominent east-west (zonal) jet streams and banded cloud patterns on Jupiter, Saturn, Uranus, and Neptune remains a major unsolved problem in planetary science. On Jupiter and Saturn, there exist $\\sim20$--30 jets, including a broad, fast superrotating (eastward) equatorial jet. In contrast, Uranus and Neptune exhibit only $\\sim3$ jets each, with high-latitude eastward jets and a subrotating (westward) equatorial jet. All four planets also exhibit a variety of compact vortices, waves, turbulent filamentary regions, short-lived convective events, and other local features. Recent observational studies demonstrate that small eddies pump momentum up-gradient into the zonal jets on Jupiter and Saturn \\citep{2006Icar..185..430S, Del_Genio2007}, which strongly suggests that cloud-layer processes are important in jet formation, although this does not exclude a possible role for the deep interior too. Scenarios to explain these diverse observations range from the ``shallow-forcing'' scenario, in which jet generation occurs via injection of turbulence, absorption of solar radiation and latent heat release in the cloud layer, to the ``deep-forcing'' scenario, in which the jet formation results from convection occurring throughout the molecular envelope \\citep[for reviews see][]{ingersoll-etal-2004, 2005RPPh...68.1935V}. Over the past several decades, a variety of idealized models have been developed that successfully produce banded zonal flows reminiscent of those on the giant planets \\citep{1978JAtS...35.1399W, cho-polvani-1996a, 1998JAtS...55..611H, Heimpel2007, lian-showman-2008}.  Despite these successes, the equatorial jet direction and magnitude have proved to be formidable puzzles that are difficult to explain. Many published models predict the same equatorial jet direction for all four giant planets and thereby fail to provide a coherent explanation that encompasses both the gas giants (Jupiter/Saturn) and the ice giants (Uranus/Neptune).  Under relevant conditions, one-layer shallow-water-type models generally produce westward equatorial flow for all four planets \\citep{1996PhFl....8.1531C,1999PhFl...11.1272I, showman-2007, scott-polvani-2007}, consistent with Uranus and Neptune but not Jupiter and Saturn.  In the parameter regime of giant planets, some shallow-water models have produced equatorial superrotation \\citep{scott-polvani-2008}, but as yet these models make no predictions for why superrotation should occur on Jupiter and Saturn but not Uranus and Neptune.  Some recent three-dimensional (3D) shallow-atmosphere models can also produce equatorial superrotation under specific conditions \\citep{Williams_2006,Williams_2002, Williams_2003_3, Williams_2003_2, Williams_2003_1, 2005P&SS...53..508Y, lian-showman-2008}, but this has generally required the addition of {\\it ad hoc} forcing, and even if such forcing were plausible, it is unclear why it would occur on Jupiter/Saturn but not Uranus/Neptune. In contrast, the deep convection models produce equatorial superrotation in most cases \\citep{2001GeoRL..28.2557A, 2001GeoRL..28.2553C, 2005Natur.438..193H,Heimpel2007, glatzmaier-etal-2008}, consistent with Jupiter and Saturn but inconsistent with Uranus and Neptune. In the context of deep convection models, \\citet{aurnou-etal-2007} proposed that Uranus and Neptune are in a regime where geostrophy breaks down in the interior, leading to turbulent mixing of angular momentum and a westward equatorial jet.  However, this mechanism occurs only at heat fluxes greatly exceeding than those observed on Uranus and Neptune \\citep{aurnou-etal-2007}.  Moreover, given that the heat fluxes on Jupiter and Saturn exceed those on Uranus and Neptune, one might expect the mechanism to apply more readily to the former pair than the latter pair; if so, one should see equatorial subrotation on Jupiter/Saturn yet superrotation on Uranus/Neptune, backward from the observed equatorial jet directions. \\citet{schneider-liu-2009} developed a 3D numerical model that produced banded zonal jets and an equatorial superrotation on Jupiter.  This is the first model that combines the deep convection (via a simple convective adjustment scheme) and absorption of solar radiation in a shallow atmosphere.  However, their model extends to only 3 bars pressure and thus neglects the effects of latent heating, which may be crucial in generating horizontal temperature contrasts in the cloud layer. Furthermore, it is unclear whether their model can produce equatorial subrotation on Uranus and Neptune by the same mechanism.  It is fair to say that we presently lack a coherent explanation for the equatorial jets that encompasses both the gas giants (Jupiter/Saturn) and the ice giants (Uranus/Neptune).  Here, we test the hypothesis that large-scale latent heating associated with condensation of water vapor can pump the zonal jets on the four giant planets. This hypothesis has been repeatedly suggested over the past 40 years \\citep{barcilon-gierasch-1970, 1976Icar...29..445G, 2000Natur.403..630I, 2000Natur.403..628G}, but this idea has not yet been adequately tested in numerical models.  While observations cannot yet constrain the existence of large-scale latent heating, abundant evidence nevertheless exists for moist convection on the giant planets. Lightning has been identified in nightside images of Jupiter from Voyager, Galileo, Cassini, and New Horizons; these flashes typically occur within localized, opaque clouds that grow to diameters up to $\\sim3000\\,$km over a few days, indicating convective activity.  The lightning illuminates finite regions on the cloud deck, indicating that the flashes occur at depths of 5--10 bars, in the expected water condensation region \\citep{borucki-williams-1986, dyudina-etal-2002}. Near such storms, clouds are sometimes observed whose tops are deeper than 4 bars, where the only condensate is water \\citep{1998Icar..135..230B, 2000Natur.403..628G}. On Saturn, electrostatic discharges presumably caused by lightning have been identified, as have explosive convective clouds that probably cause them \\citep{Porco2005, dyudina-etal-2007}. Whistlers and electrostatic discharges indicating the presence of lightning have also been detected on Uranus and Neptune \\citep{zarka-pedersen-1986, gurnett-etal-1990, kaiser-etal-1991}, suggesting that moist convection occurs on these planets too.  For plausible water abundances (a few times solar or greater), latent heating can cause local temperature increases great enough to have important meteorological effects.  To date, numerical models of jet formation have generally not included moisture and its latent heat release. Most two-dimensional (2D) and shallow-water models adopt  forcing that injects turbulence everywhere simultaneously and is confined to a small range of wavenumbers \\citep[e.g.][] {1998JAtS...55..611H, scott-polvani-2007}, which does not capture the sporadic and localized nature of moist-convective events.   Likewise, existing 3D models of Jovian jet formation have been dry (no water vapor) and force the flow by imposing latitudinal temperature differences rather than including moist convection \\citep[e.g.][]{Williams_2006,Williams_2002, Williams_2003_3, Williams_2003_2, Williams_2003_1, 2005P&SS...53..508Y, lian-showman-2008}. Notable efforts in the right direction are the one-layer studies by \\citet{Li-2006} and \\citet{showman-2007}, which adopted a forcing explicitly intended to represent the effects of moist convection. \\citet{Li-2006} adopted a quasigeostrophic model and introduced isolated vorticity patches to represent moist-convective storms; \\citet{showman-2007} adopted the shallow-water equations and introduced isolated mass pulses to represent the moist convection.   These studies show that, under planetary rotation, the small-scale turbulent flow can inverse cascade to form large scale dynamics:  zonal jets dominate at low latitudes and vortices dominate at high latitudes.   Nevertheless,  these models do not explicitly include water vapor, and the moist convection events are, rather than occurring naturally, injected by hand with prescribed sizes, lifetimes and amplitudes.   Studies have also been carried out that investigate the effects of sophisticated cloud microphysics schemes on the vertical structure in 1D column models \\citep{Del_Genio1990} and 2D height/latitude models \\citep{nakajima-etal-2000, 2008Icar..194..303P}, but these studies do not address whether moist convection can pump the jets.  Our previous 3-D studies with imposed latitudinal temperature variation can produce baroclinic eddies that drive the zonal jets through inverse-cascade of turbulence \\citep{lian-showman-2008}. Those simulations successfully reproduced some major dynamic features on Jupiter such as banded zonal winds and equatorial superrotation; they also predicted that the jets on Jupiter could extend significantly deeper than the eddy accelerations that pump them.   However, the nature of the imposed forcing schemes was only a crude parameterization of the processes that produce latitudinal temperature contrasts.  Here we present three-dimensional (3D) global numerical simulations using the MITgcm to investigate whether large-scale latent heating can drive the zonal jets on Jupiter, Saturn, Uranus, and Neptune.  Specifically, we investigate whether we can explain (i) the approximate number and speed of jets, and (ii) the direction of the equatorial jet on all four planets in the context of a single mechanism. We explicitly include water vapor as a tracer in our numerical model. Section 2 describes the numerical model, section 3 presents the simulation results, and section 4 concludes. ",
        "conclusions": "We presented global, three-dimensional numerical models to simulate the formation of zonal jets by large-scale latent heating on the giant planets. These models explicitly include water vapor and its condensation and the resulting latent heating. We find that latent heating can naturally produce banded zonal jets similar to those observed on the giant planets. Our Jupiter and Saturn simulations develop $\\rm \\sim 20$ zonal jets and Uranus/Neptune simulations develop $\\rm 3 - 7$ zonal jets depending on the water abundance. The jet spacing is consistent with Rhines scale $\\pi (2U/\\beta)^{1/2}$. The zonal jets in our simulations produce modest violations of the barotropic and Charney-Stern stability criteria at some latitudes.  In our simulations, condensation of water vapor releases latent heat and produces baroclinic eddies.  These eddies interact with the large-scale flow and the $\\beta$ effect to pump momentum up-gradient into the zonal jets.  At the same time, a meridional circulation develops whose Coriolis acceleration counteracts the eddy accelerations in the weather layer. This near-cancellation of eddy and Coriolis accelerations leads to slow evolution of the zonal jets and maintains the steadiness of zonal jets in the presence of continual forcing. Such a process was suggested by \\citet{2006Icar..182..513S} and \\citet{Del_Genio2007} and occurred also in the simulations of \\citet{showman-2007} and \\citet{lian-showman-2008}.  Our simulations also produce equatorial superrotation for Jupiter and Saturn and subrotation for Uranus and Neptune.  Although a number of previous attempts have been made to produce superrotation on Jupiter/Saturn and subrotation on Uranus/Neptune, previous models have generally lacked an ability to produce {\\it both} superrotation on Jupiter/Saturn and subrotation on Uranus/Neptune without introducing {\\it ad hoc} forcing or tuning of model parameters. Ours is the first study to naturally produce such dual behavior, without tuning, within the context of a single model. Although the speeds of the superrotation and subrotation are weaker than observed, our simulations provide a possible mechanism to explain the dichotomy in equatorial-jet direction between the gas giants (Jupiter/Saturn) and ice giants (Uranus/Neptune) within the context of a single model. In our simulations, the strength, scale, and direction of the equatorial flow are strongly affected by the abundance of water vapor as well as by the planetary radius and rotation rate. Equatorial superrotation preferably forms in simulations with modest water vapor abundance and is further promoted by large planetary radii and fast rotation rates, as occur at Jupiter and Saturn.  In contrast, high water abundance leads to equatorial subrotation regardless of the planetary radius and rotation rate explored here.  In this picture, the dichotomy in the equatorial jet direction between the gas and ice giants would result from a combination of the faster rotation rates, larger radii, and probable lower water abundances on Jupiter/Saturn relative to those on Uranus/Neptune.  Despite this encouraging result, Saturn poses a possible difficulty with this picture.  Our Saturn simulations generated equatorial superrotation when the water abundance is five times solar but equatorial subrotation when it is ten times solar.  Saturn's actual water abundance is unknown but could easily lie anywhere within this range.  Moreover, even our five-times-solar Saturn case produced a superrotation that is much weaker and narrower than observed --- about $120\\rm \\,m\\,s^{-1}$ in the simulation versus $\\sim400\\rm \\,m\\,s^{-1}$ on the real planet.  (The equatorial jet in our Jupiter simulations is also too narrow, although the discrepancy is less severe.)  However, a variety of processes excluded here (e.g., cloud microphysics, realistic radiative transfer, and moist convection) could significantly affect the jet profile.  Future investigations that include these effects are needed before a definite assessment can be made.  Consistent with our earlier work, we again find that our simulations generate deep barotropic jets despite the localization of the eddy accelerations to the weather layer (pressures less than $\\sim7\\,$bars in our Jupiter simulations, for example). In some simulations, the deep jets have similar strength as the winds in weather layer. The mechanism that forms the deep jets is similar to the mechanism we previously identified \\citep{lian-showman-2008,2006Icar..182..513S}. However, these deep jets are affected by the redistribution of water vapor, which makes the diagnostics very complicated in the present case.  Our simulations successfully produce the large-scale dynamic features on Jupiter and Uranus/Neptune under the effect of large-scale latent heating. However, our simulations lack long-lived vortices such as the Great Red Spot on Jupiter and Great Dark Spot on Neptune. We also ignore the precipitation and re-evaporation of condensates and use an idealized radiative cooling scheme. The grid resolution in our simulations is relatively low for resolving the mesoscale moist convection events which have typical horizontal scales of $1000$ kilometers \\citep{Little1999, Porco2003}. Future models can include cloud physics, a sub-grid-scale parameterization of cumulus convection, a more realistic radiative transfer scheme, and explore the coupling between the deep interior and the weather-layer processes identified here."
    },
    "0910/0910.4159_arXiv.txt": {
        "abstract": "\\noindent We study large scale structure in the cosmology of Coleman-de Luccia bubble collisions. Within a set of controlled approximations we calculate the effects on galaxy motion seen from inside a bubble which has undergone such a collision. We find that generically bubble collisions lead to a coherent bulk flow of galaxies on some part of our sky, the details of which depend on the initial conditions of the collision and redshift to the galaxy in question. With other parameters held fixed the effects weaken as the amount of inflation inside our bubble grows, but can produce measurable flows past the number of efolds required to solve the flatness and horizon problems. ",
        "introduction": "One of the greatest challenges for string theory is to find a testable prediction which can differentiate it from other potential theories of quantum gravity and high energy physics. Unfortunately, the typical energy scales at which string theoretic effects become important is likely far out of reach for any earth-based particle accelerator. Cosmology though, provides another opportunity. The early universe was arbitrarily hot and dense and signatures of the physics at that time remain in the large-scale structure of the universe today.   It has become increasingly clear in recent years that one of the characteristic features of string theory is the ``landscape'' model of cosmology \\cite{discretum,kklt,lennylandscape}.  There are many ways to generate realistic four-dimensional universes in string theory via flux compactifications. Various parameters specifying the compactification become scalar fields in our four-dimensional universe. The landscape is a potential energy function of the would-be moduli scalar fields, which possesses many distinct minima.  These minima are meta-stable and can decay via bubble nucleation-type tunneling or quantum thermal processes.  In the early universe, one expects all the minima to be populated as the temperature decreases and the fields begin to settle.  The minima with the highest values of potential energy then begin to inflate, and are expected to rapidly dominate the volume.  After a few string times the universe will be almost completely full of rapidly inflating regions, and as time continues to pass the metastable field values in these regions will begin to decay by bubble nucleation.  Since we live in a region with very small vacuum energy, we expect that our observable universe is contained in such a bubble.  The most promising way to test this idea is to look at large scale cosmological structure.  If we model the string theory landscape as a scalar field coupled to gravity with a potential with several minima, universes inside a bubble which formed via a tunneling instanton have certain characteristic features:  they are open, they are curvature dominated at the ``big bang\" (rendering it non-singular), and they have a characteristic spectrum of density perturbations generated during the inflationary period that takes place inside them  \\cite{cdl,gott1,gott2,Garriga:1998he,lindeopen,Freivogel:2005vv,batra}.   Moreover, because the bubble forms inside a metastable region, other bubbles will appear in the space around it.  If these bubbles appear close enough to ours, they will collide with it \\cite{oldguth,ggv,Aguirre:2007an}.  The two bubbles then ``stick\" together, separated by a domain wall, and the force of the collision releases a pulse of energy which propagates through each of the pair.  This process (as seen from inside our bubble) makes the universe anisotropic, affects the cosmic microwave background and the formation of structure.  A collision between two bubbles produces a spacetime which can be solved for exactly \\cite{ben,wwc,Aguirre:2007wm,wwc2}.  In \\cite{wwc,wwc2}, one of these authors began the study of this process and its effects. In these papers we principally studied the effects of such a collision on the cosmic microwave background (CMB).  We found that bubble collisions naturally create hemispherical power asymmetries spread through the lower CMB multipoles and/or produce relatively small cold or hot spots, depending on the time of the collision (and hence the size of the disk), and can lead to measurable non-Gaussianities.  The qualitative features, however, are always the same:  a collision with a single other bubble produces anisotropies that depend only on one angle---the angle to the vector pointing from our location to the center of the other bubble---and which are generically monotonic in the angle.  While the CMB temperature map is a good probe of these effects, it is by no means the only one. In this paper, we begin to look for other measurable effects by studying how a single bubble collision affects large scale structure and galaxy motion. We find that collisions generically lead to a large scale coherent flow of galaxies in our sky. The size of the affected disk of galaxies, and the magnitude of the effect are controlled by similar order parameters for the CMB effect, as well as the redshifts of the galaxies in question. Generically, the angular size of the affected disk on the sky for large scale structure is always larger than the corresponding disk for the CMB hot/cold spots. It is very interesting to consider our results in light of recent observational measurements of bulk galaxy flows dubbed ``dark flow'' \\cite{darkflow,Kashlinsky:2008us,Watkins:2008hf}.  Our model consists of a collision between our bubble -- which we treat as a thin-wall Coleman-de Luccia bubble which undergoes a period of inflation after it formed (which we model as a de Sitter phase) followed by a period of radiation domination -- and another different bubble, the details of which are largely unimportant.  We assume that the tension of the domain wall between the bubbles and the vacuum energy in the other bubble are such that the domain wall accelerates away from us.  Such collisions can produce observable effects which are not in conflict with current data.  We will also assume that the inflaton is the field that underwent the initial tunneling transition, so that the reheating surface is determined by the evolution of that field and can be affected by the collision with the other bubble through its non-linear equation of motion, and that no other scalar fields are relevant.  After some number of efolds of inflation, our bubble will reheat and become radiation dominated. The subsequent expansion and the geometry of the reheating surface itself will be modified by the presence of the domain wall. In \\cite{wwc2}, it was shown how the surface was modified. This modified surface will also backreact on the spacetime because the various fluids created at reheating will no longer be classically comoving. In \\cite{wwc2} this backreaction was ignored and  it was assumed the universe could be approximated as standard radiation dominated Friedmann-Robertson-Walker (FRW) universe for the subsequent evolution. We will show here that this approximation is valid for the CMB photons, but not for the matter fluid that will form the galaxies we are interested in here. To find the evolution of galaxies, some measure of this backreaction must be taken into account, and we make a first approximate attempt to do so here.   There are a number of anomalies in current cosmological data, including a cold spot \\cite{coldspot1, coldspot2, coldspot3, coldspot4, coldspot5}, hemispherical power asymmetry \\cite{Eriksen:2007pc}, non-Gaussianities (see {\\it e.g.} \\cite{nongauss}),  the ``axis of evil\" \\cite{Land:2005ad}, and the tentative possibility of a ``dark flow\" \\cite{darkflow,Kashlinsky:2008us,Watkins:2008hf}.  While the significance of these effects is unclear, there are few if any first-principle models which can account for them.  Our results here, combined with \\cite{wwc,wwc2} indicate that bubble collisions can naturally create hemispherical power asymmetries spread through the lower CMB multipoles, relatively small cold or hot spots, measurable non-Gaussianities and a possible macroscopic, coherent bulk flow of galaxies. While it is not yet clear that a single collision can produce all of these effects in agreement with observations, it is nevertheless interesting that bubble collisions can at least qualitatively produce many of these effects. ",
        "conclusions": "In this paper we have shown that bubble collisions generically lead to a large scale bulk flow of galaxies. There is a disk on the sky of affected galaxies, always larger than the corresponding spot for the CMB temperature map, that decreases in size with redshift (though it may take up the entire sky for all redshifts observed). Within this disk, galaxies will coherently flow, with a peculiar velocity that decreases in magnitude from the center of the disk and is proportional to the corresponding temperature at the center of the spot in the CMB map. This is a distinctive signal profile and is in principle measurable with improved observations or analysis of current cluster catalogs. It would be very interesting if a bulk galaxy flow could be correlated with a corresponding disk or spot in the CMB data as this would be evidence for bubble collisions and the string landscape.  Recently, two groups have claimed to observe bulk flows of galaxy clusters \\cite{Watkins:2008hf,darkflow,Kashlinsky:2008us}. They both measure a flow of $\\mathcal{O}(400-800 \\, \\trm{km/s})$ out to a redshift $z \\sim 0.2$. Both measurements are subject to a variety of random and systematic uncertainties (e.g. filtering methods and the conversion from temperature to velocity). They seem to have reasonable directional agreement, but the flow measured by Watkins et. al is considerably smaller than that of Kashlinsky et al. and they measure over different redshift ranges (Watkins et. al catalog is dominated by objects with $z<0.03$ and Kashlinsky by objects with $0.04<z<0.2$). The flows we find appear to be an order of magnitude smaller, as our flows are constrained by the corresponding effect in the CMB. Larger flows would lead to a very cold or hot spot in the CMB, likely ruled out by current observations. In addition, since our peculiar velocity profile falls off the further we are observing from the center of the disk of affected galaxies, and the disk itself may not take up the whole sky, the corresponding average peculiar velocity measured would be smaller. However, we have made several approximations which could affect the measured flow: \\begin{itemize} \\item{We have ignored any self-interaction and gravitation of the radiation and matter fluids. It is possible that as the fluids clump, this could affect the backreacted metric in region III and hence the flow.} \\item{We have ignored the transition to matter domination. Taking this transition into account would likely increase the flow for a subtle reason. In a matter dominated universe peculiar velocities will redshift more quickly. This implies that galaxies entering region III will transition more quickly to flowing along $\\partial_{\\tau_{III}}$. Transforming back to the CMB rest frame this results in higher {\\it measured} peculiar velocity with respect to the rest frame.} \\item{The small freedom we have in choosing the metric interpolation in region III can have an $\\mathcal{O}(1)$ effect on the flow. The flows computed above were for the choice that seems to produce the smallest flow. Other choices increase the flow velocity.} \\item{We have likely not accounted for many more detailed effects of the cosmological evolution that could affect the flow.} \\end{itemize} Because of these approximations, it is not hard to envision a more detailed treatment finding an enhanced flow.  Many open questions on bubble collisions remain. It would be interesting to carry out a more rigorous model of cosmological evolution and perturbation theory in this background and find a more accurate picture of these effects. In this paper and \\cite{wwc,wwc2} we have mapped out two possible effects of bubble collisions resulting from the string landscape, namely the effect on the CMB temperature map and on galaxy cluster motion. One could search for other effects, e.g. polarization of the CMB as our toy models bear some relation to those of \\cite{Dvorkin:2007jp} and it may be possible to further correlate these signals. This effect could be particularly interesting in light of new polarization data expected from the Planck experiment. One could also look for signals in 21cm radiation data. On the observational side, most of the current measurements of the CMB have been presented in momentum ($\\ell$) space, yet these effects are most apparent in real space. Each collision would lead to a disk on the sky and one could search the CMB and other data for traces of disks, some of which might overlap. While many of these things are likely to be difficult or require improved measurements (e.g. 21cm telescopes), we feel that cosmology and astrophysics offers a unique opportunity to possibly detect effects from high energy physics and string theory."
    },
    "0910/0910.1898_arXiv.txt": {
        "abstract": "We investigate the young stellar population in and near the cometary globule Ori\\,I-2. The analysis is based on deep Nordic Optical Telescope $R$-band and H$\\alpha$ images,  JCMT SCUBA 450 and 850~\\micron\\ images combined with near-infrared 2MASS photometry and mid-infrared archival {\\em Spitzer} images obtained with the IRAC (3.6, 4.5, 5.8 and 8~\\micron), and MIPS (24 and 70~\\micron) instruments. We identify a total of 125 sources within the 5\\arcmin$\\times$5\\arcmin\\ region imaged by IRAC. Of these sources 87 are detected in the $R$-band image and 51 are detected in the 2MASS survey. The detailed physical properties of the sources are explored using a combination of near/mid-infrared color-color diagrams, greybody fitting of SEDs and an online SED fitting tool that uses a library of 2D radiation transfer based accretion models of young stellar objects with disks. Ori\\,I-2 shows clear evidence of triggered star formation with four young low luminosity pre-main sequence stars embedded in the globule. At least two, possibly as many as four, additional low-mass PMS objects, were discovered in the field which are probably part of the young $\\sigma$ Orionis cluster.  Among the PMS stars which have formed in the globule, MIR-54 is a young, deeply embedded Class 0/I object, MIR-51 and 52 are young Class II sources, while MIR-89 is a more evolved, heavily extincted Class II object with its apparent colors mimicking a Class 0/I object.  The Class 0/I  object MIR-54 coincides with a previously known IRAS source and is a strong sub-millimeter source. It is most likely the source for the molecular outflow and the large parsec scale Herbig-Haro flow. However the nearby Class II source, MIR-52, which is strong a H$\\alpha$ emission line star, also appears to drive an outflow approximately aligned with the outflow from MIR-54, and because of the proximity of the two outflows, either star could contribute.  MIR-89 appears to excite a low excitation HH object, HH~992, discovered for the first time in this study. ",
        "introduction": "Cometary globules are believed to be molecular cloud condensations, which are so dense that they are not disrupted when an \\ion{H}{2} region expands into  the molecular cloud(s) surrounding it.  Such globules are always bright rimmed and often have a cometary or tear shaped form, because the stellar wind and ionizing radiation from the central O stars ablate away the low density gas on the side facing the O stars and sweep away dust towards the tail. Cometary globules are thought to form from ``elephant trunks''  \\citep{Herbig74,Reipurth84}, although for compact cores surrounded by low density gas this phase can be quite short. The Rosette nebula is a good example of an \\ion{H}{2} region, which shows an abundance of both cometary globules and elephant trunk structures \\citep{Herbig74}. Although elephant trunks and cometary globules have been known for a long time, the fact that these sites may be forming stars was realized only in the last few decades.  The first clear confirmation that stars form in cometary globules was a result of the identification of Bernes~135 in the cometary globule CG\\,1, one of the large cometary globules in the Gum Nebula, as a pre-main-sequence (PMS) star \\citep{Reipurth83}.  The discovery of molecular outflows associated with cold IRAS sources without optical counterparts in three bright-rimmed globules \\citep{sugitani1989}, definitely confirmed that stars form in elephant trunks and cometary globules. Probably the most well-known example of star forming in the tip of elephant trunks is seen in the Hubble poster,  ``the pillars of creation'',  in the Eagle nebula (M\\,16) \\citep{Hester96}.   In this paper we examine star formation  in a more nearby cometary globule, Ori I-2, located in the large \\ion{H}{2} region IC\\,434. This is one of the three globules, in which \\citet{sugitani1989} discovered a molecular outflow.  By using {\\it Spitzer} IRAC and MIPS images, 2MASS data, deep R and H$\\alpha$ images from the Nordic Optical Telescope (NOT), and SCUBA sub-millimeter imaging we can identify all young stars in and near the globule and provide more accurate information on their physical characteristics. ",
        "conclusions": "We find clear evidence for triggered star formation in the cometary globule Ori\\,I-2, with four young stars embedded in the front side of the globule. The PMS object furthest away from the bright rim and approximately at the center of the globule,  MIR-54, is a young Class I or stage I protostar.  It is most likely responsible for the large parsec scale Herbig-Haro outflow, with its counter flow seen as a bright H$\\alpha$ outflow breaking out on the western side of the globule. However, MIR-52, the heavily obscured H$\\alpha$ emission line star $\\sim$ 20\\arcsec\\ south of MIR-54, also appears to drive an outflow, which is approximately aligned with the outflow from MIR-54.  To the east both outflows overlap, and either star could therefore be the exciting star for the large HH outflow. MIR-89, which is located in the bright rim, is a more evolved Class II object, which  most likely excites the HH object, HH\\,992, which we discovered in this study. HH\\,992 lies outside the globule, $\\sim$ 65\\arcsec\\ east of the star and south of the large Herbig-Haro outflow,  HH\\,289.  HH\\,992 breaks up into three knots, and must be rather low excitation, since it is not seen in the H$\\alpha$ image.  Another candidate young star, MIR-104, which is close to the globule and near MIR-89 (Figure~\\ref{fig_halpha_rband}) may also have formed in the globule.  However, MIR-104 was only detected at 5.8 and 8 \\micron\\ and we therefore have insufficient information to determine whether it really is a PMS object. If it is, it could  have formed inside the globule before the gas was ablated away by the strong UV radiation from $\\sigma$ Ori.  We identify two infrared excess stars, one of which is far to the south, $\\sim$2\\arcmin\\  from the globule and the other is seen in projection against the bright rim in the tail region of the globule. Both these sources have zero or very little foreground extinction and  are probably unrelated to the globule, although we cannot exclude the possibility that they could have formed in the globule in an earlier epoch of star formation. We also find two other candidate PMS stars in the field, which may have weak mid-IR excesses. All these stars appear older, i.e. they are pre-transitional or transitional disk stars, which now are part of $\\sigma$-Orionis cluster, which has an age of $\\sim$ 3 Myr.  Models of cometary globules can successfully explain their morphology and suggest that they form stars through Radiation Driven Implosion (RDI)   \\citep{lefloch1994,miao2006}.  This was already qualitatively discussed by \\citet{Reipurth83}, who showed that cometary globules are active sites of star formation. In these models cometary globules are denser cores in the molecular cloud surrounding an \\ion{H}{2} region, which get exposed to the UV radiation from the central OB association, when the \\ion{H}{2} region expands into it. The UV flux of the OB association ionizes the external layers of the core facing the OB association and an ionization front is formed at the front surface.  The gas is heated during the ionization and the temperature increases. The increased pressure due to the hot ionized gas drives an isothermal shock into the cloud, which compresses the neutral gas in the core. At the same time, the ionized and heated gas at the surface of the core flows radially away from the surface and forms an  evaporating layer surrounding the surface of the core. Due to the curvature of the front surface of the core the shock waves preceding the ionization front converge and collide at their focus, causing the interior of core to be compressed to much higher densities. This causes the central part of the core to become gravitationally unstable, triggering the formation of a star.  The observations of \\citet{ikeda2008} of the isolated cometary globule BRC\\,37 in the \\ion{H}{2} region IC\\,1396 appears to match this model rather well. Their observations show a string of PMS objects along the symmetry axis of the globule, with the oldest being in front of the globule, and the youngest object still embedded in the head of the globule, suggesting sequential formation of young stars as the ionization front proceeds towards the tail of the globule. For some of the  other cometary globules in the same \\ion{H}{2} region, like IC\\,1396\\,N (BRC\\,38), this scenario is less clear. \\citet{getman2007} found an elongated clustering of X-ray sources aligned with the globule, with the youngest stars still embedded inside it, consistent with the RDI model for triggered star formation. However \\citet{beltran2009}, who studied the same cometary globule using deep J, H, K$^\\prime$, and narrowband H$_2$ 2.12~\\micron\\ imaging and previously published mm imaging did not find any color or age gradient in the south-north direction, i.e. in the direction of the globule. They therefore conclude that not all star formation in this globule can be explained in terms of triggering. We can think of two possible explanations for why IC\\,1396\\,N does not appear to show the same type of age progression and sequential triggering as seen in BRC\\,37. If the X-ray sources in front of the globule are late Class II objects, i.e. transitional disk like objects, or Class III objects, they will show no color excess in the near-IR and therefore appear as field stars. Secondly, IC\\,1396\\,N is not a well defined isolated cometary globule with a simple head/tail structure, suggesting that it could consist of more than one cometary globule.  Ori\\,I-2, the cometary globule studied in this paper, does appear more similar to BRC\\,38. It shows clear evidence for triggered star formation, with the youngest object, MIR-54, being furthest away from the front side of the globule. The two other young Class II objects are definitely older and closer to the bright rim and the ionizing star, $\\sigma$ Orionis. All  three stars are also close to the symmetry axis of the globule, as predicted by the RDI model. MIR-89 and possibly MIR-104, which both lie on the eastern side of the globule, do not appear to have formed by  focussed compression along the symmetry axis of the globule. Other processes, like Rayleigh-Taylor instabilities in the bright rim, may also play a role and form low-mass stars like MIR-89. A similar situation is seen in the Horsehead, a bright rimmed elephant trunk structure east of $\\sigma$ Orionis, which also shows some evidence for sequential star formation triggered by the expanding \\ion{H}{2} region \\citep{mookerjea2009}. These young PMS stars are also low-mass stars, which appear to have formed in the bright rim by Rayleigh-Taylor instabilities, since hardly any of them lie close to the symmetry axis of the Horsehead.  Several, if not all of the young stars in Ori\\,I-2 drive outflows. It is interesting to note that all of the outflows are roughly perpendicular to the symmetry axis of the globule. It maybe a coincidence, but other cometary globules with embedded protostars, like BRC\\,37, show a similar behavior. If this is generally true, it would provide additional support to the focussed RDI model, which predicts that the gas inside a cometary globule will be compressed in a ridge along the symmetry axis of the globule. When density condensations, i.e. cores in this ridge become gravitationally unstable and collapse, the collapse would preferentially be towards the symmetry axis, therefore aligning the protoplanetary disk along the symmetry axis. Since outflows are driven by disks, one would therefore see outflows which are perpendicular to the symmetry axis of the cometary globule.   Ori\\,I-2 is thus a low mass star forming cometary globule. The star formation in this region appears to be triggered by a combination of RDI and Rayleigh-Taylor instabilities in the rim of the globule. In the deep R-band images for the first time we detect HH~992, which based on its non-detection in the \\halpha\\ image is most likely an extremely low excitation HH object.  Follow-up optical and infrared spectroscopic observations to understand the true nature of HH~992 as well as MIR-52, the sole \\halpha\\ emission star in the field, are needed.  We find that many of the young stars drive outflows which are aligned perpendicular to the symmetry axis of the globule. Higher spatial resolution imaging of the outflows in Ori\\,I-2 in near-infrared and millimeter would provide structural details of the outflows and lead to correct identification of the sources exciting these outflows."
    },
    "0910/0910.2643_arXiv.txt": {
        "abstract": "{% The formation and evolution of galaxies is imprinted on their stellar population radial gradients.  Two recent articles present conflicting results concerning the mass dependence of the metallicity gradients for early-type dwarf galaxies. On one side, Spolaor et al.  show a tight positive correlation between the total metallicity \\zh{} and the mass.  On the other side, in a distinct sample\\thanks{Based on observations made with ESO telescopes at La Silla Paranal observatory under program ID076.B-0196}, we do not find any trend involving \\feh{} (Koleva et al.). In order to investigate the origin of the discrepancy, we examine various factors that may affect the determination of the gradients: namely the sky subtraction and the signal-to-noise ratio. We conclude that our detection of gradients are well above the possible analysis biases. Then, we measured the \\mgf{} relative abundance profile and found moderate gradients.  The derived \\zh{} gradients scatter around -0.4\\,dex/$r_e$.  The two samples contain the same types of objects and the reason of the disagreement is still not understood.  } ",
        "introduction": "The stellar populations of galaxies  hold a fossil record of their formation and  evolution.  Their  study may allow  to discriminate between  the  various  formation  scenarios  predicting  different spatial  gradients and different  relations between  the gradients and the mass of the galaxy.  In  the  classical  monolithic  scenario,  galaxies  formed  in  a dissipative collapse \\citep{lar1974,ay1987}, where stars remain on their orbits and do not mix.  The gas, enriched in metals from the evolved  stars,   flows  into  the   centre,  generating  negative metallicity  gradients (\\emph{i.e.}  higher  metal content  in the centre  than in the  outskirts).  The  chemo-dynamical simulations \\citeauthor{mat1987}  \\citep[1987,  see also][]{kaw2003}  predict, that due to  their small potential, dwarf galaxies,  are unable to retain  metals and thus  their gradients  are shallower  or almost inexistent.   Therefore  relations  between   the  mass   and  the metallicity or the metallicity gradient are expected.  A   competing  (or  complementary)   formation  scenario   is  the hierarchical  clustering \\citep[e.g.][]{col1994}, where  small dark matter halos  merge to produce bigg\\-er ga\\-la\\-xi\\-es.  It may be thought that  these violent events mix the  stellar population and erase the  gradients \\citep[e.g.][]{whi1980}, but  some simulations \\citep[e.g.][]{dim2009}, at variance, show a gradient preservation, due to the violent relaxation which does not mix the orbits.  The   interpretation   of  the   metallicity   gradients  is   not straightforward.   Nevertheless,   gradients   observation   could constrain model's  predictions. Several studies on  big samples of giant   elliptical   galaxies    have   been   already   performed \\citep[e.g.][]{psb2006}.   However,   due   to  their   low-surface brightness,  the  dwarf  galaxies,  are  still  very  much  'terra incognita'.  Recently we published  a paper on the stellar  content of 16 dwarf galaxies in the Fornax  cluster and in groups \\citep{kol2009}. One of  our main  conclusions was  that most  of our  galaxies  hold a strong negative metallicity  gradient.  The star formation history analyses  revealed  that  these   gradients  are  built  at  early epochs. We did not find  any correlation between our gradients and the  velocity  dispersion  (proxy  of  the  mass)  in  our  sample (Fig.\\ref{fig:gradients}).  However, this  result is challenged by a  recent letter from  \\citet{spolaor2009}, using  a sample  of 14 dwarfs in the Virgo and  Fornax clusters observed with GMOS at the Gemini  South telescope in  medium resolution,  long-slit, optical spectroscopy.  To analyse  their data, \\citet{spolaor2009} binned  them along the slit  direction  to achieve  a  minimum  signal-to-noise  of 30  @ 5200\\,\\AA.   They measured  Lick indices  \\citep{wor1994,wo1997} and derived  ages  and  chemical  compositions  by  comparing  to  the \\citet{tho2004}  single  stellar  population (SSP)  models,  using their  own method \\citep{pro2002}.   They estimated  the gradients performing a  linear least-squares  fit to the  radial metallicity profile weighted by their errors.  The fit was done in logarithmic space  of the  radius, $\\Delta$\\zh/$\\Delta\\log(r/r_e)$,  where $r$ varies from 1 arcsec to one effective radius ($r_e$).   \\begin{figure}[ht!] \\includegraphics{koleva_fig01a.eps} \\includegraphics{koleva_fig01b.eps} \\caption[]{   Metallicity  gradients  \\emph{vs.}   $\\sigma$  and \\magb{}  (top and  bottom respectively)  of the  2  samples of dwarf  elliptical  galaxies,   one  from  \\citet[][blue,  open squares]{spolaor2009}  and   other  is  this   work  (filled symbols).  Green  circles are galaxies  from \\citet{vaku}, red circles   from   \\citet{kol2009}    and   a   red   star   for NGC\\,205\\citep{sp2002}.   The  names of  our  galaxies are  also plotted.} \\label{fig:gradients} \\end{figure}  There  is no  galaxy in  common between  the two  samples  and the causes of the disagreement may  be multiple: choice of the sample, data reduction  or analysis issue.   The primary purpose  of the present  paper is  to explore  these possibilities.   In  order to directly compare our results  with \\citet{spolaor2009} we will determine our metallicity gradients in an identical way.  In Sect.2  we present the  sample, observations and  analysis.  In Sect.\\ref{section:snr}    we    investigate     the    systematics (signal-to-noise ratio, bad  sky-subtraction) which may affect the derived  metallicity gradients.  As  \\citeauthor{spolaor2009} used the total metallicity \\zh we  also measured \\mgf{} and discuss the differences   between    $\\Delta$\\feh{}   and   $\\Delta$\\zh{}   in Sect.\\ref{subsec:feh}. ",
        "conclusions": "We investigated the  discrepancy in the metallicity gradients-mass relation  between  two recent  works,  namely \\citet{spolaor2009} and \\citet{kol2009}. While the former  work shows a very strong and tight relation between the  total metallicity gradient in low-mass galaxies and their central  velocity dispersion or \\magb, the latter do not find a hint of such a relation.  As first guess  for this discrepancy we pointed  the small size of the samples and different  selection criteria.  Such a scenario is possible,    since    some     of    our    galaxies    fall    on \\citeauthor{spolaor2009}   relation.   This   hypothesis  may   be supported   by   the   fact    that   the   Fornax   galaxies   of \\citet{spolaor2009} are in  general flat ($\\epsilon\\,>\\,0.3$), and we  have shown  that the  gradients in  such galaxies  are usually shallower.  However, in their Virgo sample all but one galaxy have $\\epsilon\\,<\\,0.3$\\footnote{ We  used the ellipticity measurements from      the      HyperLeda      database      \\citep{pat2003}, \\url{http://leda.univ-lyon1.fr}}.  Hence,  this  explanation  is ruled out and flattening cannot be a reason for our disagreement.  Another possible explanation would  have been that while the dwarf galaxies exhibit a strong  negative $\\Delta$\\feh, they possess the same \\afe{}  gradient but with  a positive sign.   Thus, measuring the  total metallicity  as a  combination of  \\feh{} and  \\afe, we would  have a  total  $\\Delta$\\zh\\,=\\,0.  We  measured the  \\mgf{} gradients  in our sample  using unpublished  semi-empirical models \\citep{php2007}.   These models are  preliminary but  the relative measurements  of \\mgf{}  are certainly  reliable. Opposite  to the expectations  we   found  that  (most  of)   our  galaxies possess \\emph{negative}   \\mgf{}  gradients.    Consequently,   the  total metallicity   gradients  are   \\emph{stronger}  than   the  \\feh{} gradients.  Negative iron and  \\mgf{} gradients will imply that,  in the early epochs,  when  the  feedback  is  a  strong  regulator,  the  star formation   starts  in   the  centre,   where  the   gas   is  the densest. After the gas is exhausted, the star formation stops first in the centre and continues  longer at larger radius. This process forming a  negative \\mgf{} gradient.  To form  the negative \\feh{} gradient  we need  that cool,  enriched gas  flows again  into the centre,  triggering  star   formation.   In  other  words,  having repeating  episodes  of  star  formation.   This  scenario  is  in agreement with  recent SPH simulations  on the formation  of dwarf galaxies \\citep[fig.10]{val2008}  We also investigated  the effect of a bad  sky-subtraction and low signal-to-noise  on our metallicity  measurements. We  showed that our metallicity measurements are not biased down to S/N\\,=\\,10 and that to have zero  metallicity gradients we need to under-subtract the sky with more than  50\\,\\%.  We concluded that these two effects cannot be the cause of our strong gradients.  As a conclusion,  according to the tests performed  in this paper, the  origin of the  differences between  the two  articles remains unknown."
    },
    "0910/0910.5914_arXiv.txt": {
        "abstract": "Axions in the $\\mu$eV mass range are a plausible cold dark matter candidate and may be detected by their conversion into microwave photons in a resonant cavity immersed in a static magnetic field. The first result from such an axion search using a superconducting first-stage amplifier (SQUID) is reported. The SQUID amplifier, replacing a conventional GaAs field-effect transistor amplifier, successfully reached axion-photon coupling sensitivity in the band set by present axion models and sets the stage for a definitive axion search utilizing near quantum-limited SQUID amplifiers. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.2705_arXiv.txt": {
        "abstract": "We present a new near-infrared survey covering the $2\\deg^2$ COSMOS field conducted using WIRCam at the Canada--France--Hawaii Telescope. By combining our near-infrared data with Subaru $B$ and $z$ images, we construct a deep, wide-field optical-infrared catalogue. At $K_{\\rm s}<23$ (AB magnitudes), our survey completeness is greater than $90\\%$ and $70\\%$ for stars and galaxies, respectively, and contains $143,466$ galaxies and $13,254$ stars. Using the $BzK$ diagram, we divide our galaxy catalogue into quiescent and star-forming galaxy candidates. At $z\\sim2$, our catalogues contain $3931$ quiescent and $25,757$ star-forming galaxies representing the largest and most secure sample at these depths and redshifts to date. Our counts of quiescent galaxies turns over at $K_{\\rm s}\\sim22$, an effect which we demonstrate cannot be due to sample incompleteness. Both the number of faint and bright quiescent objects in our catalogues exceeds the predictions of a recent semi-analytic model of galaxy formation, indicating potentially the need for further refinements in the amount of merging and active galactic nucleus feedback at $z\\sim2$ in these models. We measure the angular correlation function for each sample and find that the slope of the field galaxy correlation function flattens to $1.5$ by $\\ks\\sim23$. At small angular scales, the angular correlation function for passive $BzK$ galaxies is considerably in excess of the clustering of dark matter.  We use precise 30-band photometric redshifts to derive the spatial correlation length and the redshift distributions for each object class. At $K_{\\rm s}<22$ we find $r_0^{\\gamma/1.8}=7.0\\pm0.5h^{-1}$~Mpc for the passive $BzK$ candidates and $4.7\\pm0.8h^{-1}$~Mpc for the star-forming $BzK$ galaxies. Our $pBzK$ galaxies have an average photometric redshift of $z_p\\sim1.4$, in approximate agreement with the limited spectroscopic information currently available.The stacked $K_{\\rm s}$ image will be made publicly available from IRSA. ",
        "introduction": "Understanding the formation and evolution of galaxies is one of the central themes in observational cosmology. To address this issue considerable work is underway to construct complete galaxy samples over a broad redshift range using many diverse ground and space-based facilities.  The ultimate aim of these studies is to construct distribution functions of galaxy properties such as stellar mass, star formation rate, morphology and nuclear activity as a function of redshift and environment and to use these measurements in conjunction with theoretical models to establish the evolutionary links from one cosmic epoch to another. Essentially we would like to understand how the galaxy population at distant times became the galaxies we see around us today and to identify the physical mechanisms driving this transformation. Improvements in both modelling and observations are the ultimate drivers in achieving this understanding.   The cold dark matter (CDM) model of structure formation remains our best description of how structures grow on large scales from minute fluctuations in the cosmic microwave background \\citep{1996MNRAS.282..347M,Springel:2006p8432}. In this picture, structures grow ``hierarchically'', with the smallest objects forming first. Unfortunately, CDM only tells us about the underlying collisionless dark component of the universe; what we observe are luminous galaxies and stars. At large scales, the relationship between dark matter and luminous objects (usually codified as the ``bias'') is simple, but at small scales (less than a few megaparsecs) the effects of baryonic physics intervene to make the relationship between luminous and non-luminous components particularly complex. Understanding how luminous objects ``light up'' in dense haloes of dark matter one is essentially the problem of understanding galaxy formation. ``Semi-analytic'' models avoid computationally intensive hydrodynamic simulations by utilising a set of scaling relations which connect dark matter to luminous objects (as will be explained later in this paper), and these models are now capable of predicting how host of galaxy properties, mass assembly and star-formation rate evolve as a function of redshift.   Observationally, at redshifts less than one, various spectroscopic surveys such as the VVDS \\citep{2005A&A...439..845L,2005A&A...439..863I,Pozzetti:2007p100} DEEP2 \\citep{Faber:2007p5801,Noeske:2007p5802} and zCSOSMOS-bright \\citep{Lilly:2007p1799,Silverman:2009p5804,Mignoli:2009p5805} have mapped the evolution of galaxy and active galactic nuclei (AGNs) populations over fairly wide areas on the sky. There is now general agreement that star formation in the Universe peaks at $1<z<2$ and that $\\sim 50\\%-70\\%$ of mass assembly took place in the redshift range $1<z<3$ \\citep{Connolly:1997p8767,Dickinson:2003p8724,Arnouts:2007p3665,Pozzetti:2007p100,Noeske:2007p5802,PerezGonzalez:2008p8736}. Alternatively stated, half of today's stellar mass appears to be in place by $z\\sim1$ \\citep{Drory:2005p8438,Fontana:2004p8883}. This is largely at odds with the predictions of hierarchical structure formation models which have difficulty in accounting for the large number of evolved systems at relatively early times in cosmic history \\citep{Fontana:2006p3228}. Furthermore, there is some evidence that around half the stellar mass in evolved or ``passive'' galaxies assembled relatively recently \\citep{Bell:2004p8521}.  It is thus of paramount importance to gather the largest sample of galaxies possible at this redshift range.   In the redshift range $1.4 < z < 3.0 $ identifiable spectral features move out of the optical wave bands and so near-infrared imaging and spectroscopy become essential. The role of environment and large-scale structure at these redshifts is largely unexplored \\citep{Renzini:2009p5799}. It is also worth mentioning that in addition to making it possible to select galaxies in this important range, near-infrared galaxy samples offer several advantages compared to purely optical selections (see, for example \\cite{Cowie:1994p1284}). They allow us to select $z>1$ galaxies in the rest-frame optical, correspond more closely to a stellar-mass-selected sample and are less prone to dust extinction. As $k-$ corrections in $K-$ band are insensitive to galaxy type over a wide redshift range, near-infrared-selected samples provide a fairly unbiased census of galaxy populations at high redshifts (providing that the extinction is not too high, as in the case of some submillimeter galaxies). Such samples represent the ideal input catalogues from which to extract targets for spectroscopic surveys as well as for determining accurate photometric redshifts.  \\cite{Cowie:1996p8471} carried out one of the first extremely deep, complete $K-$ selected surveys and made the important discovery that star-forming galaxies at low redshifts have smaller masses than actively star-forming galaxies at $z\\sim1$, a phenomenon known as ``downsizing''. Stated another way, the sites of star-formation ``migrate'' from higher-mass systems at high redshift lower-mass systems at lower redshifts.  More recent $K$-selected surveys include the K20 survey \\citep{Cimatti:2002p3013} reaching $\\kab\\simeq 21.8$ and the GDDS survey \\citep{Abraham:2004p2980} which reached $\\kab\\simeq 22.4$ provide further evidence for this picture. The areas covered by these surveys was small, comprising only $\\sim 55$ arcmin$^2$ and $\\sim 30$ arcmin$^2$ K20 and GDDS respectively. While \\cite{Glazebrook:2004p3320} and \\cite{Cimatti:2008p1800} provided spectroscopic confirmation of evolved systems $z>1.4$ and provided further evidence for the downsizing picture \\citep{Juneau:2005p8688} their limited coverage made them highly susceptible to the effects of cosmic variance. It became increasingly clear that much larger samples of passively evolving galaxies were necessary.   At $K<20$ the number of passive galaxies at $z\\sim2$ redshifts is small and spectroscopic followup of a complete magnitude-limited sample can be time-consuming. For this reason a number of groups have proposed and validated techniques based on applying cuts in colour-colour space to isolate populations in certain redshift ranges. Starting with the Lyman-break selection at $z\\sim3$ \\citep{Steidel:1996p9040}, similar techniques have been applied at intermediate redshifts to select extremely red objects (EROs; \\cite{Hu:1994p9049}) or distant red galaxies (DRGs; \\cite{Franx:2003p310}) and the ``BzK'' technique used in this paper \\citep{Daddi:2004p76}. The advantage of these methods is that they are easy to apply requiring at most only three or four photometric bands; their disadvantage being that the relationships between each object class is complicated and some selection classes contain galaxies with a broad range of intrinsic properties \\citep{Daddi:2004p76,Lane:2007p295,Grazian:2007p398}. The relationship to the underlying complete galaxy population can also be difficult to interpret \\citep{LeFevre:2005p8609}. Ideally, one like to make complete mass-selected samples at a range of redshifts but such calculations require coverage in many wave bands and can depend sensitively on the template set \\citep{Pozzetti:2007p100,Longhetti:2009p9068}. Moreover, for redder populations the mass uncertainties can be even larger; \\cite{Conroy:2009p9107} estimate errors as larger as 0.6 dex at $z\\sim 2$.   At $z\\sim1.4$ \\cite{Daddi:2004p76} used spectroscopic data from the K20 survey in combination with stellar evolutionary tracks to define their ``BzK'' technique.  They demonstrated that in the $(B-z)$ $(z-K)$ colour-colour plane, star-forming galaxies and evolved systems are well separated at $z>1.4$, making it possible accumulate larger samples of passive galaxies at intermediate redshifts that was possible previously with simple one-colour criterion.   Subsequently, several other surveys have applied these techniques to larger samples of near-infrared selected galaxies. In one of the widest surveys to date, \\cite{Kong:2006p294} constructed $K$-band selected samples over a $\\sim 920$ arcmin$^2$ field reaching $\\kab\\simeq 20.8$ reaching to $\\kab\\simeq 21.8$ over a 320 arcmin$^2$ sub-field. The exploration of a field of this size made possible to measure the clustering properties of star-forming and passive galaxy sample and to establish that passively evolving galaxies in this redshift range are substantially more strongly clustered than star-forming ones, indicating that a galaxy-type - density relation reminiscent of the local morphology-density relation must be already in place at $z\\gsim 1.4$.   The UKIDSS survey reaches $\\kab\\sim22.5$ over a $\\sim 0.62$-deg$^2$ area included in the Subaru-\\textit{XMM Newton} Deep Survey and \\cite{Lane:2007p295} used this data set to investigate the different commonly-used selected techniques at intermediate redshifts, concluding most bright DRG galaxies have spectra energy distributions consistent with dusty star-forming galaxies or AGNs at $\\sim2$. They observe a turn-over in the number counts of passive $BzK$ galaxies.   Other recent works include the MUSYC/ECDFS survey covering $\\sim 900$ arcmin$^2$ to $\\kab\\sim 22.5$ over the CDF South field \\citep{Taylor:2008p3500}, not to be confused with the GOODS-MUSIC catalog of \\cite{Fontana:2006p3228}, which covers 160 arcmin$^2$ of GOODS-South field to $\\kab\\sim 23.8$. This $K$-band selected catalogue, as well as the FIREWORKS catalog by \\cite{Wuyts:2008p5806}, are based on the ESO Imaging Survey coverage of the GOODS-South field\\footnote{http://www.eso.org/science/goods/releases/20050930./} These studies have investigated, amongst other topics, the evolution of the mass function at $z\\sim2$ and what number of red sequence galaxies which were already in place at $z\\sim2$.   Finally, one should mention that measuring the distribution of a ``tracer'' population, either red passive galaxies or normal field galaxies can provide useful additional information on the galaxy formation process. In particular one can estimate the mass of the dark matter haloes hosting the tracer population and, given a suitable model for halo evolution, identify the present-day descendants of the tracer population, as has been done for Lyman break galaxies at $z\\sim3$. A few studies have attempted this for passive galaxies at $z\\sim2$, but small fields of view have made these studies somewhat sensitive to the effects of cosmic variance.  The ``COSMOS'' project \\citep{2007ApJS..172....1S} comprising a contiguous $2~\\deg^2$ equatorial field with extensive multi-wavelength coverage, is well suited to probing the universe at intermediate redshift.   In this paper we describe a $K_{\\rm s}$-band survey covering the entire $\\sim 1.9\\deg^2$ COSMOS field carried out with WIRCam at the Canada-France-Hawaii Telescope (CFHT). The addition of deep, high resolution $K-$ data to the COSMOS field enables us to address many of scientific issues outlined in this introduction, in particular to address the nature of the massive galaxy population in the redshift range $1<z<2$.  Our principal aims in this paper are to (1) present a catalogue of $BzK$-selected galaxies in the COSMOS field; (2) present the number counts and clustering properties of this sample in order principally to establish the catalogue reliability; (3) present the COSMOS $K-$ imaging data for the benefit of other papers which make indirect use of this dataset (for example, in the computation of photometric redshifts and stellar masses).  Several papers in preparation or in press make use of the data presented here. Notably, \\cite{Ilbert:2009p7351} combine this data with IRAC and optical data to investigate the evolution of the galaxy mass function. The deep part of the zCOSMOS survey \\citep{Lilly:2007p5770} is currently collecting large numbers galaxies spectra at $z>1$ in the central part of the COSMOS field using a colour-selection based on the $K$- band data set described here.   Throughout the paper we use a flat lambda cosmology ($\\Omega_m~=~0.3$, $\\Omega_\\Lambda~=~0.7$) with $h~=~H_{\\rm0}/100$~km~s$^{-1}$~Mpc$^{-1}$. All magnitudes are given in the AB system, unless otherwise stated. The stacked $K_{\\rm s}$ image presented in this paper will made publicly available at IRSA.\\footnote{\\url{http://irsa.ipac.caltech.edu/Missions/cosmos.html}} ",
        "conclusions": "\\label{sec:summary}  We have presented counts, colours and clustering properties for a large sample of $K-$ selected galaxies in the $2\\deg^2$ COSMOS-WIRCam survey. This represents the largest sample of galaxies to date at this magnitude limit. By adding deep Subaru $B-$ and $z-$ data we are able to classify our catalogue into star-forming and quiescent/passive objects using the selection criterion proposed by \\cite{Daddi:2004p76}. To $K_{\\rm s}<23.0$ our catalogues comprises $143,466$ galaxies, of which $3931$ are classified as passive galaxies and $25,757$ as star-forming galaxies. We have also identified a large sample of $13,254$ faint stars.  Counts of field galaxies and star-forming galaxies change slope at $K_{\\rm s}\\sim22$. Our number counts of quiescent galaxies turns over at $\\ks\\sim22$, confirming an observation previously made in shallower surveys \\citep{Lane:2007p295}. This effect cannot be explained by incompleteness in any of our very deep optical bands. Our number counts of passive, star-forming and field galaxies agree well with surveys with brighter magnitude limits.  We have compared our counts to objects selected in a semi-analytic model of galaxy formation. For simple magnitude-limited samples the \\cite{Kitzbichler:2007p3449} model reproduces very well galaxy counts in the range $16<\\ks<20$. However, at fainter magnitudes \\citeauthor{Kitzbichler:2007p3449}'s model predict many more objects than are observed.  Comparing this model with predictions of passive galaxy counts, we find that at $20<\\kab<20.5$ model counts are below observations by a factor of 2, whereas at $22.5<\\kab<23.0$ model counts are in excess of observations by around a factor of 1.5. This implies that the \\citeauthor{Kitzbichler:2007p3449} model predicts too many small, low-mass passively-evolving galaxies and too few large high-mass passively evolving galaxies at $z\\sim1.4$. In these models, bulge formation takes place by mergers. At $\\ks\\sim 22$, passive galaxies in the millennium simulation have stellar masses of $\\sim 10^{11}M_\\odot$, similar to spectroscopic measurements of passive galaxies \\citep{Kriek:2006p7151}. This suggests that the difference between models and observations is linked to the amount of ``late merging'' taking place \\citep{DeLucia:2007p7510}. The exact choice of the AGN feedback model can also sensitively affect the amount star-formation in massive systems \\citep{DeLucia:2006p9608,Bower:2006p7511}. It is clear that observations of the abundance of massive galaxies can now provide insight into physical processes occurring in galaxies at intermediate redshifts. For the time being it remains a challenge for these models to reproduce both these observations at high redshift and lower-redshift reference samples.  Our results complement determinations of the galaxy stellar mass function at intermediate redshifts which show that total mass in stars formed in semi-analytic models is too low at $z\\sim2$ compared to models \\citep{Fontana:2006p3228}. We note that convolution with standard uncertainties of $\\sim0.25$ dex in mass function estimates at $z\\sim2$ can make a significant difference in the mass function, as can be see in Figure~14 of \\cite{Wang:2008p9639} who show detailed comparisons between semi-analytic models and observations. The discrepancy between our observations and models cannot be explained in this way.  We have cross-matched our catalogue with precise 30-band photometric redshifts calculated by \\citeauthor{Ilbert:2009p4457} and have used this to derive the redshift distributions for each galaxy population. At \\kab$\\sim 22$ our passive galaxies have a redshift distribution with $z_{\\rm {med}}\\sim1.4$, in approximate agreement with similar spectroscopic surveys comprising smaller numbers of objects. Most of our $pBzK$ galaxies have $z_p<2.0$, in contrast with the redshift distribution for $sBzK$ galaxies and for the general field galaxy population which extend to much higher redshifts at this magnitude limit. In the redshift range $1<z<3$ at $\\ks\\sim22$ the $pBzK$ population represents around $\\sim20\\%$ of the total number of galaxies, in contrast to $\\sim80\\%$ for $sBzK$-selected galaxies. DRG-selected galaxies remain an important fraction of the total galaxy population, reaching around $\\sim 50\\%$ of the total at $z\\sim2$.  Our work confirms that most galaxies satisfying the passive-$BzK$ selection criteria lie in a narrower redshift than either $sBzK$- or DRG-selected objects. Interestingly, a few a passive $BzK$ galaxies in our survey have $z_p>2.0$, and it is tempting to associate these objects with higher-redshift evolved galaxies detected in spectroscopic surveys \\citep{Kriek:2008p9702}. We  have investigated the clustering properties of our catalogues for which the $2\\deg^2$ field of view of the COSMOS survey provides a unique probe of the distant universe. Our stellar correlation function is zero at all angular scales to $K_{\\rm s}\\sim23$ demonstrating the photometric homogeneity and stability of our catalogues. For a $K_{\\rm s}-$ selected samples, the clustering amplitude declines monotonically toward fainter magnitudes. However, the slope of the best-fitting angular correlation function becomes progressively shallower at fainter magnitudes, an effect already seen in the COSMOS optical catalogues.  At the faintest magnitude slices, the field galaxy population (all objects with $18.0<\\ks<23.0$) is only slightly more clustered than the underlying dark matter distribution, indicating that $K_{\\rm s}-$ selected samples are excellent tracers of the underlying mass. On the other hand, star-forming and passive galaxy candidates are more clustered than the field galaxy population. At arcminute scales and smaller the passive BzK population is strongly biased with respect to the dark matter distribution with bias values of 2.5 and higher, depending on scale.  Using our photometric redshift distributions, we have derived the comoving correlation length $r_0$ for each galaxy class. Fitting simultaneously for slope and amplitude we find a comoving correlation length $r_0^{\\gamma/1.8}$of $\\sim7 h^{-1}$~Mpc for the passive $BzK$ population and $\\sim 5 h^{-1}$~Mpc for the star-forming $BzK$ galaxies at $K_{\\rm s}<22$ . Our field galaxy clustering amplitudes are in approximate agreement with optically-selected red galaxies at lower redshifts.  High bias values are consistent with a picture in which $pBzK$ galaxies inhabit relatively massive dark matter haloes on order of $\\sim10^{12}$M$_\\odot$, compared to the $sBzK$ and field galaxy population. We will return to this point in future papers, where will interpret these measurements in terms of the halo model.  Measuring spectroscopic redshifts for a major fraction of 3000 $pBzK$ galaxies in the COSMOS field one will have to wait for the advent of large-throughput, wide-field near-infrared ($J$-band) spectrographs on 8-10m class telescopes, such as FMOS at Subaru (Kimura et al. 2006). Smaller field, cryogenic, multi-object spectrographs such as MOIRCS at Subaru \\citep{Ichikawa:2006p4968}, EMIR at the GTC telescope \\citep{Garzon:2006p4837}, and Lucifer at the LBT \\citep{Mandel:2006p4846} should prove effective in producing high S/N spectra for a relatively small fraction of the $pBzK$ galaxies in the COSMOS field.  Since this paper was prepared, additional COSMOS-WIRCam $K_{\\rm s}$ data observations have been taken which will increase the total exposure time by $\\sim30\\%$. In addition, new $H$- observations have also been made. Both these data products will be made publicly available in around one year from the publication of this article. In the longer term, the COSMOS field will be observed as part of the UltraVISTA deep near-infrared survey which will provide extremely deep $JHK$ observations over the central part of the field."
    },
    "0910/0910.3229_arXiv.txt": {
        "abstract": "The 18806 \\rosat\\ All Sky Survey Bright Source Catalog (\\bsc) \\xray\\ sources are quantitatively cross-associated with near-infrared (\\ir) sources from the Two Micron All Sky Survey Point Source Catalog (\\psc). An association catalog is presented, listing the most likely counterpart for each \\bsc\\ source, the probability $P_{\\rm id}$ that the \\ir\\ source and \\xray\\ source are uniquely associated, and the probability $P_{\\rm no-id}$ that none of the \\psc\\ sources are associated with the \\xray\\ source. The catalog includes 3853 high quality ($P_{\\rm id}>0.98$) \\xray--\\ir\\ matches, 2280 medium quality ($0.98\\geq P_{\\rm id}>0.9$) matches, and 4153 low quality ($0.9\\geq P_{\\rm id}>0.5$) matches. Of the high quality matches, 1418 are associations that are not listed in the \\simbad\\ database, and for which no high quality match with a \\usno\\ optical source was presented for the \\bsc\\ source in previous work. The present work offers a significant number of new associations with \\bsc\\ objects that will require optical/\\ir\\ spectroscopy for classification.  For example, of the 6133 $P_{\\rm id}>0.9$ \\psc\\ counterparts presented in the association catalog, 2411 have no classification listed in the \\simbad\\ database. These \\psc\\ sources will likely include scientifically useful examples of known source classes of \\xray\\ emitters (white dwarfs, coronally active stars, active galactic nuclei), but may also contain previously unknown source classes. It is determined that all coronally active stars in the \\bsc\\ should have a counterpart in the \\psc, and that the unique association of these \\bsc\\ sources with their \\ir\\ counterparts thus is confusion limited. ",
        "introduction": "\\label{sec:intro}  The \\rosat\\ All Sky Survey Bright Source Catalog \\citep[\\bsc ;][]{voges99} contains flux and positional information for 18806 X-ray sources\\footnote{The reference \\citep{voges99} refers to 18811 sources, but 5 duplicates were removed in the most recent version of the catalog. See \\url{ftp://ftp.xray.mpe.mpg.de/rosat/catalogues/rass-bsc/1.3\\_vs\\_1.4.diff} for details (accessed 2009 May 5).}, down to a limiting 0.1--2.4 keV count rate of $0.05\\ {\\rm counts\\, s}^{-1}$.  The observations were made in 1990/91 with the \\rosat/PSPC satellite, and at a brightness limit of $0.1\\ {\\rm counts\\, s}^{-1}$ (5843 sources) the \\bsc\\ represents a sky coverage of 92\\%. The median positional uncertainty ($1\\sigma$, including 6\\arcsec\\ systematic uncertainty) is $11\\arcsec$.  The sources in the \\bsc\\ are the brightest \\xray\\ sources in the sky. Brighter sources generally allow for better \\xray\\ spectra and time variability observations, so the \\bsc\\ sources are ideal for studying the \\xray\\ properties of their respective classes with modern \\xray\\ telescopes. However, as of 2008 Aug 1 only 1982 \\bsc\\ sources have been observed with either {\\em Chandra} or {\\em XMM-Newton}. Further, around two thirds of the sources in the \\bsc\\ remain unclassified (\\cf\\ \\S\\ref{sec:prevclass}), so the number of known source classes in the \\bsc\\ and the number of sources of each class is limited.  Classifying \\bsc\\ sources therefore not only expands our knowledge of the brightest \\xray\\ sources in the sky, but also provides ideal targets for follow-up observations that can improve our understanding of specific source classes.  As a larger fraction of the \\bsc\\ is classified, the catalog will also become a useful tool in producing population constraints for various classes of objects, and the sources that remain unclassified can be targeted in searches for previously unknown source classes and objects with unusual properties.  To classify \\bsc\\ sources, one needs spectral and flux measurements in other wavebands. The error cones of the \\xray\\ sources often contain several off-band sources, any one (or none) of which may be a counterpart. If the probability of association between the \\xray\\ source and one or more off-band counterparts is calculated, the classification of that \\xray--off-band pair/group can be constrained. Furthermore, probable off-band counterparts can be studied with ground-based or other instrumentation, which can provide constraints on the source of X-rays. Similarly, the absence of off-band sources near a \\bsc\\ source can also constrain its class.  Thus, searching for off-band counterparts to \\bsc\\ sources is a good way to advance the classification and study of these sources.  A summary of previous work identifying likely counterparts for ${>500}$ \\bsc\\ sources is given in Table \\ref{table:cross-associations}, which also includes the present work for comparison. Many of these previous cross-associations searched for counterparts in catalogs of objects of a specific type, like stars \\citep{huensch98, huensch99, makarov00, torres06}, galaxies \\citep{zimmermann01}, or quasars \\citep{bade98, zickgraf03}.  Other previous cross-associations searched for counterparts in catalogs of sources in one or more specific wave-bands, like optical \\citep{voges99, rutledge00, schwope00, mcglynn04} and/or radio \\citep{bauer00, mcglynn04}.  Only work that specifically addresses finding counterparts to \\bsc\\ sources has been included in Table \\ref{table:cross-associations}, while work that includes other \\xray\\ catalogs without explicitly considering the \\bsc\\ subset \\citep[e.g.][]{flesch04} has not been included. \\placetable{table:cross-associations}  The \\bsc\\ sources were quantitatively cross-associated with \\usno\\ optical sources in previous work \\citep{rutledge00}. The present work adapts the methods used in that \\xray--optical cross-association to associate \\bsc\\ sources with Two Micron All Sky Survey Point Source Catalog \\citep[\\psc ;][]{skrutskie06} near-infrared (\\ir) sources. The present work is the first to produce a cross-association between the full \\bsc\\ and \\psc\\ catalogs.  \\twomass\\ surveyed 99.998\\% of the celestial sphere in the near-infrared, resulting in a Point Source Catalog containing 470992970 point sources, and an Extended Source Catalog (XSC) containing 1647599 extended sources. The majority of \\xsc\\ sources are galaxies ($\\sim$97\\%), while most \\psc\\ sources are stars in the Milky Way. The \\psc\\ also contains point-source processed versions of virtually all the \\xsc\\ sources, as well as a significant number of unresolved galaxies \\citep{cutri03}. The present work considers only the \\psc, for which the greatest sensitivity was achieved in the J band ($1.12-1.36\\mu m$), with a $10\\sigma$ detection level better than 15.8 mag \\citep{skrutskie06}.  This paper is organized as follows. In \\S\\ref{sec:method}, the method used to perform the statistical cross-association between \\bsc\\ and \\psc\\ sources is described. The results of this cross-association are discussed in \\S\\ref{sec:results}, which includes an association catalog containing the most likely \\ir\\ counterpart for each \\xray\\ source. In \\S\\ref{sec:prevclass}, the source types listed in the \\simbad\\ database for the \\psc\\ sources that appear in the association catalog are discussed. The present work is compared to previous cross-associations in \\S\\ref{sec:comparison}, and discussion and conclusions can be found in \\S\\ref{sec:conclude}. ",
        "conclusions": "\\label{sec:conclude}  The association catalog presented in \\S\\ref{subsec:catalog} lists each \\bsc\\ source, the corresponding $P_{\\rm no-id}$, the most likely \\psc\\ counterpart, and the $P_{\\rm id}$ of the \\bsc--\\psc\\ match. The catalog includes 3853 high quality matches ($P_{\\rm id}>0.98$), of which $\\leq39$ are expected to be chance associations with background objects. Of the high quality matches, only 1921 are with a \\bsc\\ source for which a $P_{\\rm id}>0.98$ \\usno\\ counterpart has previously been identified \\xidp. Thus, between the present work and the \\bsc--\\usno\\ cross-association, there are 4637 \\bsc\\ sources with a $P_{\\rm id}>0.98$ optical and/or \\ir\\ counterpart. The present work thus represents a significant improvement over previous cross-associations.  The purpose of uniquely associating \\bsc\\ and \\psc\\ sources is to obtain J-, H-, and K-band fluxes for, and better localization of, \\bsc\\ sources, to aid in classification.  Certain source classes appear to be restricted to specific ranges of \\xray--J color and J--K color (\\cf\\ \\S\\ref{subsec:color}), so the \\xray, J-, and K-band fluxes of matches can be used to select interesting objects for precise classification through follow-up observations.  Previous classification of the proposed \\psc\\ counterpart in a match can also provide \\bsc\\ classifications.  The present work presents likely new classifications for 1425 previously unclassified \\bsc\\ sources, based on the classification of the $P_{\\rm id}>0.98$ \\psc\\ counterpart.  All the coronally active stars in the \\bsc\\ should have counterparts in the \\psc\\ (\\cf\\ \\S\\ref{subsec:color}), so unique association of these \\bsc\\ sources with \\ir\\ counterparts is confusion limited. For any coronally active stars in the \\bsc\\ that do not presently have a known \\psc\\ counterpart, better \\xray\\ localizations should allow the \\psc\\ counterpart to be found. Further, since the fraction of the \\bsc\\ sources that do not have counterparts in the \\psc\\ is $12\\pm1\\%$ (\\cf\\ \\S\\ref{subsec:fields}), there are $2200\\pm200$ \\bsc\\ sources with no \\ir\\ counterparts in the \\psc\\ catalog. These \\bsc\\ sources are likely associated with \\ir\\ sources fainter than the \\psc\\ detection limit, and represent a sizeable discovery space for classes of \\xray\\ sources other than coronally active stars.  \\bsc\\ sources with $P_{\\rm no-id}\\approx 1$ have no likely \\psc\\ counterparts.  Some of these sources may be associated with an extended near infra-red source not listed in the \\psc, some may be associated with \\ir\\ sources too far from the \\bsc\\ position to be identified, and some may be \\ir\\ faint objects. The number of \\bsc\\ sources with $P_{\\rm no-id}>0.98$ is 334, and these provide a sample which can be searched for examples of \\ir\\ faint \\xray\\ sources like isolated neutron stars."
    },
    "0910/0910.1640_arXiv.txt": {
        "abstract": "{} {Spectral line emissivities have usually been calculated for a Maxwellian electron distribution. But many theoretical works on galaxy groups and clusters and on the solar corona suggest to consider modified Maxwellian electron distribution functions to fit observed X-ray spectra. Here we examine the influence of high energy electron populations on measurements of metal abundances. } {A generalized approach which was proposed in the paper by Prokhorov et al. (2009) is used to calculate the line emissivities for a modified Maxwellian distribution. We study metal abundances in galaxy groups and clusters where hard X-ray excess emission was observed. } {We found that for modified Maxwellian distributions the argon abundance decreases for the HCG 62 group, the iron abundance decreases for the Centaurus cluster and the oxygen abundance decreases for the solar corona with respect to the case of a Maxwellian distribution. Therefore, metal abundance measurements are a promising tool to test the presence of high energy electron populations.} {} ",
        "introduction": "Galaxy clusters are large structures in the Universe, with radii of the order of a megaparsec. Groups of galaxies are the poorest class of galaxy clusters. The space between galaxies in clusters is filled with low-density $\\sim 10^{-3}$ cm$^{-3}$ high temperature ($k_{\\mathrm{B}} T\\sim 1-10$ keV) gas (for a review, e.g. Sarazin 1986). The temperatures of 1-10 keV are close to the values of the K-shell ionization potentials ($I_{\\mathrm{Z}}=Z^2 Ry$, where Z is the atomic number and Ry is the Rydberg constant) of heavy elements with atomic numbers in the range of Z=10-26.  Emission lines from heavy elements were detected by X-ray telescopes from galaxy groups and clusters. The current instruments (XMM-Newton, Chandra and Suzaku) have provided precise measurements of the chemical abundances of many elements (O, Ne, Mg, Si, C, Ar, Ca, Fe and Ni) in groups and clusters. Metal abundances around 0.5 in Solar units of Anders \\& Grevesse (1989) are derived under the assumptions of collisional equilibrium (for a review, see Werner et al. 2008).  The ionization rates, recombination rates and emissivity in a spectral line have usually been calculated for a Maxwellian electron distribution (e.g. Mewe \\& Gronenschild 1981). However, in many low-density astrophysical plasmas, the electron distribution may differ from a Maxwellian distribution (e.g. Porquet et al. 2001).  Hard X-ray tails reported from BeppoSAX observations in the X-ray spectra of some galaxy clusters (Fusco-Femiano et al. 1999; Fusco-Femiano et al. 2004; Rossetti \\& Molendi 2004 for the Coma cluster; Kaastra et al. 1999 for the Abell 2199; Molendi et al. 2002 for the Centaurus cluster) were interpreted as bremsstrahlung emission from non-thermal subrelativistic electrons (see e.g. Sarazin \\& Kempner 2000) or from thermal electrons with a Maxwellian spectrum distorted by a particle acceleration mechanism (Blasi 2000; Liang et al. 2002, Dogiel et al. 2007). The bremsstrahlung interpretation is associated with a huge energy output of emitting particles and faces energetics problems (e.g. Petrosian 2001). Evidence for a hard X-ray excess above the thermal emission was also discovered in galaxy groups with ASCA (Fukazawa et al. 2001, Nakazawa et al. 2007). The evidence for and the nature of hard X-ray spectral tails in these galaxy groups and clusters are discussed in the review by Rephaeli et al. (2008).  The use of non-extensive thermo-statistics (Tsallis 1988; for a review, see Tsallis 1999), based on the natural generalization of entropy for systems with long-range interactions, was proposed by Hansen (2005) to fit the X-ray spectrum observed near NGC 4874 (near the center of the Coma cluster). We consider non-extensive thermo-statistics as another approach to explain hard X-ray excess in groups and clusters in the framework of the bremsstrahlung model.  A more traditional interpretation of hard X-ray tails based on the inverse Compton scattering (ICS) of relativistic electrons on relic photons (Sarazin \\& Lieu 1998) faces a serious problem. The combination of hard X-ray and radio observations within the ICS model implies a magnetic field much lower than the values derived from Faraday rotation measurements (e.g. Clarke et al. 2001). Yet several arguments have been proposed to alleviate (at least in part) the disagreement (for a review, see Brunetti 2003; Ferrari et al. 2008; Petrosian et al. 2008).  The presence of high energy subrelativistic electrons (non-thermal subrelativistic electrons or thermal electrons with a Maxwellian spectrum distorted by the particle acceleration mechanism) or the use of non-extensive thermo-statistics must be probed using various observational methods in order to test interpretations of X-ray tails from galaxy clusters.  The Sunyaev-Zel'dovich (SZ) effect can be used to constrain the electron distribution in galaxy clusters. The study of the influence of high energy subrelativistic electrons on the SZ effect was done for the Coma and Abell 2199 clusters by Blasi et al. (2000) and Shimon \\& Rephaeli (2002). A method based on the measurement of the spectral slope around the crossover frequency of the SZ effect was proposed by Colafrancesco et al. (2009) to analyse the high energy electron populations in galaxy clusters.  A new probe to study the electron distribution in galaxy clusters, namely the flux ratio of the emission lines due to FeK$\\alpha$ transitions (FeXXV and FeXXVI) was considered by Prokhorov et al. (2009). This flux ratio is very sensitive to the population of electrons with energies higher than the ionization potential of a FeXXV ion (which is $\\approx$ 8.8 keV). The influence of the high energy subrelativistic electron population on the flux ratio is more prominent in low temperature clusters (as Abell 2199) than in high temperature clusters (as Coma), because the fraction of thermal electrons with energies higher than the helium-like iron ionization potential in low temperature clusters is smaller than that in high temperature clusters. However, the FeXXVI line is weak in low temperature clusters and current instruments do not have sufficient sensitivity to measure the iron line flux ratio.  Kaastra et al. (2009) have shown that relative intensities of the satellite lines are sensitive to the presence of supra-thermal electrons in galaxy clusters and that the instruments on future missions like Astro-H and IXO will be able to demonstrate the presence or absence of these supra-thermal electrons.  In this paper we study the influence of high energy electron populations on metal abundance estimates in galaxy groups and clusters and show that the effect of high energy particles can be significant. This effect is a promising test to the presence of high energy subrelativitic electrons in galaxy groups and clusters because of substantial changes in abundance estimates for modified Maxwellian distributions. We also consider the effect of high energy electrons on abundance estimates in the solar corona where the presence of modified Maxwellian electron distributions has been proposed.  The paper is organized as follows. In Sect.~2.1 we choose a galaxy group and a galaxy cluster where high energy subrelativistic electron populations have been proposed and derive values of the electron distribution parameters. We calculate the changes in metal abundances with respect to the values for a Maxwellian distribution in Sect.~2.2. We discuss the bremsstrahlung model of hard X-ray emission from galaxy clusters in Sect.~3 and present our conclusions in Sect.~4. We calculate an oxygen abundance drop for the solar corona in Appendix A. ",
        "conclusions": "We have shown in this paper that the metal abundance estimates depend on the presence of high energy subrelativistic electrons proposed to account for measurements of hard X-ray excess emission from galaxy groups and clusters. Due to the impact of these energetic electron populations, the Ar abundance estimate in the HCG 62 group and the Fe abundance estimate in the Centaurus cluster significantly decrease by $\\approx30$ and $\\approx15\\%$, respectively.  These decreases in the Ar and Fe abundance estimates are determined by the high energy subrelativistic electron fractions (6\\% for the HCG 62 group and 5\\% for the Centaurus cluster) which are comparable to the thermal electron fractions with energies higher than the K-shell ionization potentials of Ar and Fe, respectively.  The influence of the high energy subrelativistic electron populations on the abundance estimates is measurable with current instruments. Therefore this probe is more efficient for detecting the high energy subrelativistic electron populations than that based on the Doppler broadening of the spectral lines proposed by Hansen (2005), which requires very high energy resolution, or than that based on the flux ratio of the emission lines (Prokhorov et al. 2009), which requires higher sensitivity instruments.  Other possibilities to produce the change in the metal abundance estimates in galaxy clusters are the effect of resonant scattering (Gilfanov et al. 1987) and the presence of multiphase hot gas - two temperature model (Buote \\& Fabian 1998, Buote 2000).  The effect of resonant scattering causes the decrement of the FeXXV line at 6.7 keV, and, therefore, the decrement of the flux ratio of the iron lines FeXXV/FeXXVI. A decrement of the Fe abundance is then produced, as in the case of a high energy subrelativistic electron population. To separate the effects of resonant scattering and the high energy subrelativistic population influence, the SZ effect from a high energy subrelativistic population can be analyzed. Following the method of Colafrancesco et al. (2009), we calculated the value of the slope of the SZ effect in the Centaurus cluster. We obtained the value of the slope $S\\approx0.033$ for both electron distributions $f_{\\mathrm{1}}(x)$ and $f_{\\mathrm{2}}(x)$ and the value of the slope $S\\approx0.028$ for a Maxwellian spectrum. Since the slope is equal to $S\\approx4.25\\times kT/(m_{\\mathrm{e}} c^2)$ for a Maxwellian electron spectrum without a high energy subrelativistic electron population (Colafrancesco et al. 2009), the value of the slope of $S=0.033$ corresponds to that at an effective temperature $kT=4.5$ keV, which is higher than the temperature $kT=3.5$ keV observed with Beppo-SAX.  Buote (2000) analyzed the ASCA data for the HCG 62 group, the same data as analyzed by Nakazawa et al. (2007), and fitted the spectrum in the frame of a two-temperature model with temperatures $kT_{\\mathrm{1}}=0.7$ keV and $kT_{\\mathrm{2}}=1.4$ keV. Nakazawa et al. (2007) showed that the spectrum of the HCG 62 group is well reproduced by the two-temperature model of Buote (2000) in the energy range below $\\sim 4$ keV, but a fit of the full spectrum requires a thermal component with an unrealistically high temperature of $\\sim17.5$ keV.  We have shown that high energy electron populations can affect derived metal abundances in galaxy groups and clusters and in the solar corona. Therefore, metal abundances are a promising tool for an analysis of the high energy subrelativistic electron component.  \\begin{acknowledgement} I am grateful to Florence Durret, Joseph Silk, William Forman, Vladimir Dogiel and Eugene Churazov for valuable suggestions and discussions and thank the referee for very useful comments. \\end{acknowledgement}"
    },
    "0910/0910.3698_arXiv.txt": {
        "abstract": "It has been suggested that dark matter particles which scatter inelastically from detector target nuclei could explain the apparent incompatibility of the DAMA modulation signal (interpreted as evidence for particle dark matter) with the null results from CDMS-II and XENON10.  Among the predictions of inelastically interacting dark matter are a suppression of low-energy events, and a population of nuclear recoil events at higher nuclear recoil equivalent energies.  This is in stark contrast to the well-known expectation of a falling exponential spectrum for the case of elastic interactions.  We present a new analysis of XENON10 dark matter search data extending to E$_{nr}=75$~keV nuclear recoil equivalent energy.  Our results exclude a significant region of previously allowed parameter space in the model of inelastically interacting dark matter.  In particular, it is found that dark matter particle masses $m_{\\chi}\\gtrsim150$~GeV are disfavored. ",
        "introduction": "Astrophysical evidence indicates that 23\\% of the mass of the universe is in the form of non-baryonic dark matter \\cite{2008hinshaw,2007jee,2006clowe}.  A well-motivated dark matter candidate particle is found in supersymmetric extensions to the Standard Model, in which the lightest supersymmetric particle (LSP) is stable.  A cosmologically interesting relic density of Weakly Interacting Massive Particles (WIMPs) arises from rather general arguments \\cite{1996jungman}, which implies that the LSP could be dark matter.  The open question of the expected mass and cross-section of particle dark matter is being addressed by numerous direct and indirect detection experiments \\cite{gaitskell2004,baudis2006}.  The nature of the dark matter particle, and its coupling to standard model particles has been the subject of much recent theoretical investigation \\cite{2009arkanihamed, 2009cui,2009alves}.  \\subsection{The XENON10 detector} The XENON10 detector is a liquid xenon time-projection chamber with an active target mass of 13.7~kg.  It was designed to directly detect galactic dark matter particles which scatter off xenon nuclei.  Typical velocities of halo-bound dark matter particles are of order 10$^{-3}$c.  This leads to a a featureless exponential recoil energy spectra for spin-independent elastic scattering of dark matter particles on a xenon target.  In the case of a 100~GeV dark matter particle mass, the predicted elastic scattering recoil energy spectrum falls by an order of magnitude from $0-30$~keV nuclear recoil equivalent energy (keVr) \\cite{1996lewin,1996jungman}.  A particle interaction in liquid xenon creates both excited and ionized xenon atoms \\cite{1978kubota}, which react with the surrounding xenon atoms to form excimers.  The excimers relax on a scale of $10^{-8}$~s with the release of scintillation photons.  This prompt scintillation light is detected by 88 photo-multiplier tubes and is referred to as the $S1$ signal.  An external electric field ($E_d=0.73$~kV/cm) across the liquid xenon target causes a large fraction of ionized electrons to be drifted away from an interaction site.  The electrons are extracted from the liquid xenon and accelerated through a few mm of xenon gas by a stronger electric field  ($\\sim10$~kV/cm), creating a secondary scintillation signal.  This scintillation light is detected by the same photo-multiplier tubes, is proportional to the number of ionized electrons and is referred to as $S2$.  The XENON10 detector discriminates between electron recoil background and the expected nuclear recoil signal from the scattering of a dark matter particle, via the distinct ratio of proportional ($S2$) to primary ($S1$) scintillation for each type of interaction.  The XENON10 collaboration previously reported WIMP-nucleon exclusion limits for spin-independent \\cite{2008xenon10SI} and spin-dependent \\cite{2008xenon10SD} elastic scattering.   \\subsection{Inelastic dark matter} In stark contrast to elastic scattering \\cite{1996lewin,1996jungman}, a model of inelastic interactions of dark matter (iDM) predicts \\cite{2001sw,2005sw} a suppression of low-energy nuclear recoils.  The recoil energy spectrum peaks near $40$~keVr, for a 100~GeV dark matter particle incident on a xenon target (and standard halo assumptions).  As explained in \\cite{2005sw}, this feature is a result of the mass difference $\\delta$ between the proposed ground and excited states of the dark matter particle.  The value of $\\delta$ is unknown and is a free parameter in the model, subject to the physical constraint that $\\delta \\gtrsim 170$~keV would kinematically forbid any scattering from dark matter particles, given an expected particle mass $m_{\\chi}=100$~GeV and a galactic escape velocity $v_{esc}=500$~km~s$^{-1}$.  In the limit $\\delta\\rightarrow0$, the iDM model reduces to the usual elastic scattering case.  At present, the iDM model is comfortably consistent \\cite{2009chang} with all reported null results (including the recent CDMS results \\cite{2009cdms}) as well as the claimed detection from DAMA \\cite{2008dama}.  Interpretation of the XENON10 results in \\cite{2009chang} was limited by the fact that in \\cite{2008xenon10SI}, data were not published beyond 45~keVr.  Since the focus in \\cite{2008xenon10SI} was on elastic interactions, that analysis was optimized for events with recoil energies in the range $4.5-26.9$~keVr.  This paper presents a new analysis of the XENON10 dark matter search data, extending to the energy range of interest for inelastic dark matter scattering.  In doing so, we find new constraints on allowed parameter space in the iDM model. ",
        "conclusions": "A primary conclusion of this work is that the events which have previously been referred to as ``anomalous'' background have been confirmed to originate in the passive layers of LXe surrounding the target, as explained in Sec. \\ref{backgrounds}. We have shown that these background events can be targeted with high efficiency by the cuts described in Sec. \\ref{ppc}.  A corollary conclusion is that the occurrence of this class of events can be explicitly prevented by the design of the detector, rather than by software cuts.  Such a situation is realized if there is no xenon outside the active target region, or (more realistically) if the photomultipliers viewing the target volume are perfectly  blind to scintillation produced in passive areas of liquid xenon, including below the cathode grid.  We have shown that the XENON10 dark matter search data exclude previously allowed regions of the DAMA-allowed parameter space, in the model of inelastic dark matter.  Specifically, dark matter particle masses $m_{\\chi}\\gtrsim150$~GeV are disfavored.  While there are events remaining in the dark matter acceptance box in Fig. \\ref{fig1}, and while these events appear consistent with the expected spectral shape for inelastic dark matter, we do not claim a detection.  It was demonstrated in Sec. \\ref{backgrounds} that the number of remaining events is reasonably consistent with the expected number of leakage events, based on electron recoil calibration data."
    },
    "0910/0910.4137_arXiv.txt": {
        "abstract": "We have calculated the effects of large scale solar flows like the meridional circulation, giant convection cells and solar rotation on the helioseismic splitting coefficients using quasi-degenerate perturbation theory (QDPT). Our investigation reveals that the effect of poloidal flows like the large scale meridional circulation are difficult to detect in observational data of the global acoustic modes since the frequency shifts are much less than the errors. However, signatures of large scale convective flows may be detected if their amplitude is sufficiently large by looking for frequency shifts due to nearly degenerate modes coupled by convection. In this comprehensive study, we attempt to put limits on the magnitude of flow velocities in giant cells by comparing the splitting coefficients obtained from the QDPT treatment with observational data. ",
        "introduction": "Solar convection is believed to be organized in a variety of spatial and temporal scales ranging from granules (size $\\sim 1$ Mm, 0.2 hr lifetime), mesogranules (size $\\sim 10$ Mm, 3 hr lifetime), supergranules (size $\\sim 30$ Mm, 1 day lifetime) to the  giant cells (size $\\ge 100$ Mm, 1 month lifetime). The power spectra of convective velocities show distinct peaks representing granules and supergranules but no distinct features at wavenumbers representative of mesogranules or giant cells (Wang 1989; Chou et al.~1991; Straus, Deubner, \\& Fleck 1992; Straus \\& Bonaccini 1997; Hathaway et al.~2000). Numerical simulations of solar convection routinely show the existence of mesogranules and giant cells (Miesch et al.~2000; Miesch et al.~2008). In this study we will only concentrate on the theoretical frequency shifts in global surface gravity and acoustic modes, denoted f- and p-modes respectively due to interaction with large scale flows namely the giant cells, the meridional circulation and rotation. The meridional circulation, believed to play an important role in magnetic flux transport by Solar dynamo modelers (Choudhuri, Sch\\\"ussler \\& Dikpati 1995; Dikpati \\& Charbonneau 1999; Chatterjee, Nandy \\& Choudhuri 2004) was observed at the solar surface by Doppler measurements of photospheric lines (Duvall 1979; LaBonte \\& Howard 1982; Komm, Howard \\& Harvey 1993; Hathaway 1996) and later by using techniques of ring diagram and time-distance helioseismology (Giles et al.~1997; Basu, Antia \\& Tripathy 1999; Gonz\\'alez Hern\\'andez et al.~1999). These studies can only measure the flow velocity in the near surface layers. This flow having a maximum surface velocity of 30 m s$^{-1}$ is apparently poleward all the way from the equator to the poles at the Solar surface. Conservation of mass requires existence of a return flow advecting mass from the poles to the equator somewhere below the surface. The dynamo models assume that the return flow occurs near the base of the convection zone, but there is no direct evidence about where the return flow is located. It is believed that the only hope of detecting such a return flow is by analyzing the effect of the meridional circulation on the global helioseismic acoustic modes (Roth \\& Stix 2008). Their investigation using quasi degenerate perturbation theory showed that theoretical p-mode frequency shifts due to meridional circulation would be $\\sim 0.1~\\mu$Hz. Such large frequency shifts may not be expected from the observed magnitude of meridional flows. For the sake of comparison, the equatorial rotational velocity at the Solar surface is $\\sim 2000$ m s$^{-1}$ and gives rise to frequency shifts $\\sim 450\\times m$ nHz as a linear function of azimuthal order $m$. Further, since the first order contribution from meridional flow calculated using the degenerate perturbation theory vanishes, one would expect the effect to be even smaller than what is suggested by the velocity. Thus it is necessary to reexamine these calculations.  Quasi degenerate perturbation theory (QDPT) was applied to calculate shifts due to flows and asphericity in Solar acoustic frequencies by Lavely \\& Ritzwoller (1992). We have used the same formulation to calculate the p-mode frequency shifts due to differential rotation, meridional circulation and giant cells of convection. In contrast to results of Roth \\& Stix (2008), our results indicate a maximum shift of about 1 nHz for a similar meridional circulation encompassing the entire solar convection zone. Traditionally, only degenerate perturbation theory (DPT) has been used to calculate the effect of rotation on the p-modes. Woodard (1989) and Vorontsov (2007) applied QDPT to calculate mixing of eigenfunctions due to rotational coupling of p-modes and concluded that the expected line profiles of the composite modes in the Solar power spectra are not distorted significantly. In this work, we use a different approach to examine if the use of quasi degenerate perturbation theory introduces significant corrections in the frequency shifts over that obtained from degenerate perturbation theory. It is necessary to check this as inversion techniques for calculating rotation rate in the solar interior are based on degenerate perturbation theory. Helioseismic inversion for rotation is also insensitive to the N--S asymmetric component of rotation. This is because in the degenerate perturbation theory, this effect also gives zero contribution. The N--S antisymmetric component of rotation has been studied near the surface and is found to be small. In this work we wish to examine its effect to check if its contribution to frequency shift are indeed small.  The giant convective cells have always been elusive to observations at the solar surface. In the past there have been studies which have failed to detect giant cell motions (LaBonte, Howard \\& Gilman 1981; Snodgrass \\& Howard 1984; Chiang, Petro \\& Foukal 1987) as well as those hinting at their existence (Cram, Durney \\& Guenther 1983; Hathaway et al.~1996; Simon \\& Strous 1997). With the availability of Dopplergrams from SOHO/MDI, it became possible to study such long-lived and large scale features more reliably. Beck, Duvall \\& Scherrer (1998) were able to detect giant cells at the solar surface with large aspect ratio ($\\sim 4$) and velocities $\\sim 10$ m s$^{-1}$ using the MDI data. Hathaway et al.~(2000) used spherical harmonic spectra from full disk measurements to detect long-lived power at $l \\le 64$. This means that the spatial extent of the features were $\\sim 2\\pi\\Rs/\\sqrt{l(l+1)} \\ge 70$~Mm. Later Ulrich (2001) identified these long-lived features not as giant cells but higher order components of the torsional oscillation. Interestingly it was earlier proposed by Snodgrass \\& Wilson (1987) that torsional oscillations occur due to modulation of differential rotation by the Coriolis force arising from meridional motions of giant cells.  The elusiveness of the giant cell features in the surface velocities lead modelers to speculate about the subsurface nature of the giant cell motions (Latour, Toomre \\& Zahn 1981). Roth \\& Stix (1999, 2003) used QDPT to calculate the effect of giant cells on p-modes and claimed that giant cells could be found by modeling the asymmetries and line broadening in the Solar power spectrum. They claimed that finite line width of the multiplets would limit the detection of the frequency splittings to vertical velocity amplitude of 100 m s$^{-1}$ or larger. However, Roth, Howe \\& Komm (2002) claimed that giant cells could be detected with current inversion methods of global helioseismology as long as they exceed an amplitude of 10 m s$^{-1}$. Usual inversion procedure for rotation neglects giant cells and assumes that the odd splitting coefficients arise only from rotation. If there are additional contributions to these coefficients, rotation inversions would not give correct results. According to Roth, Howe \\& Komm (2002) the effect of giant cells would appear as distortions in the rotation inversion if their velocities are $\\ge 10$ m s$^{-1}$. However, using the same technique of QDPT we find that with the present global helioseismic data we can put an upper limit of 50--100 m s$^{-1}$ on the vertical velocity associated with giant cell flows. This is mainly because our calculations show much smaller effects from giant cells.  This paper is organized as follows. In \\S 2 we define zonal and poloidal flows inside the Sun, while \\S3 discusses the quasi degenerate perturbation theory, its differences from degenerate theory as well as the details of our calculations. We present our theoretical results and compare them with present f- and p-mode data from both GONG and SOHO/MDI instruments in \\S4. Finally conclusions are described in \\S5. ",
        "conclusions": "In this study we have calculated the effect of differential rotation, the meridional circulation and the giant cell flows on p-mode frequencies using quasi-degenerate perturbation theory. For toroidal flows like differential rotation, the quasi degenerate theory provides only a second order correction to the eigenfrequencies, whereas for the poloidal flows it becomes necessary to use the quasi degenerate treatment since the diagonal terms of the perturbation matrix vanishes. We agree with Lavely \\& Ritzwoller (1992) that the effect of rotation on the odd coefficients is negligible and hence using the degenerate theory is sufficient. However, the even splitting coefficients $a_{2q}$ are about one-half of the errors in observations and they might not be negligible while calculating the $a_{2q}$ for other effects like the magnetic field. This contribution is comparable to the effect of centrifugal term.  The frequency shift due to N--S antisymmetric component of rotation rate vanishes in the degenerate perturbation treatment and it is necessary to use QDPT. We find that with realistic magnitude of velocities in the antisymmetric component, the computed splitting coefficients are very small and this effect may be neglected.   For a meridional circulation with maximum horizontal velocity at the surface of $30$ m s$^{-1}$ and one cell per hemisphere we find that $l a_2 \\sim 8.5$ nHz for multiplets with turning points $r_t$ near the surface. Because of our choice of velocity field, the horizontal velocity increases rapidly with depth near the surface. This behavior is not seen in the Sun. In that case it may be more meaningful to choose velocity profile with a maximum value of 30 m s$^{-1}$, which occurs a little below the surface. In that case the splitting coefficients would need to be reduced by a factor of 4 or more, making them even smaller. The mean frequency shift due to meridional flow is found to be up to 12 nHz, which is much smaller than what is found by Roth \\& Stix (2008). The reason for this discrepancy is not clear. This could be due to differences in where the velocity is normalized to specified value.  In any case, $la_2$ is very small for multiplets with $r_t/\\Rs < 0.9$ coupled with single celled meridional circulation. This value is less than roughly one-seventh of the errors in the observed splitting coefficients. We also do not find any change with increasing the number of cells in the radial direction or by changing the depth of penetration of the meridional circulation while keeping the surface amplitude constant. This makes it impossible to comment on the return flow believed to be present at the base of the convection zone. When we increase the number of cells per hemisphere the selection rules allow more multiplets to couple with each other. In the process some of the nearly degenerate multiplets combine to give rather large values of $l a_2$. Interestingly, most of these nearly degenerate modes have turning points in the range $0.6\\Rs$ to $0.8\\Rs$. These purely meridional flows give frequency shifts which are symmetric about $m = 0$ and only give non zero values for even splitting coefficients. The observed amplitude of these higher order components near the surface is rather small and with such amplitudes the effect is expected to be small. It appears that in general the magnitude of splitting coefficients due to $u_0=9$ m s$^{-1}$ with $s=2$ is smaller than that with $u_0=100$ m s$^{-1}$ and $s=8$, even after excluding modes with nearly degenerate frequencies. This could be due to a larger maximum horizontal velocity as well as larger radial velocity (for $s=8$) or some geometric factors arising in the calculations.  In addition to this, we have also calculated the asymmetric frequency shifts caused due to giant convection cells with angular flow  profiles given by $Y_8^4(\\theta, \\phi)$ and $Y_8^8(\\theta, \\phi)$. Our results differ significantly from earlier works of Roth, Howe \\& Komm (2002) and Roth \\& Stix (2008) who find the effect of poloidal flows on p-mode frequencies to be an order of magnitude larger than what we find. The asymmetric shifts also contribute to the theoretically calculated odd coefficients. We have performed a 1.5d rotation inversion on these coefficients to detect any discernible feature in the inverted profile. However, since the magnitude of the features in the inverted profile are smaller than the inversion errors we conclude that giant cells with $u_0 \\le 100$ m s$^{-1}$ do not have much effect on the rotation inversion.  Finally we have analyzed some of the GONG data sets covering solar cycle~23 to look for possible evidence of giant cells. We do not find any signal of these flows in selected modes with large splitting coefficients and we can put upper limits on velocities of such cellular flows. From our analysis one can say that the existence of giant convection cells of $Y_8^4(\\theta, \\phi)$ angular pattern with maximum vertical velocities of 100 m s$^{-1}$  can be ruled out with a confidence level of $1.5\\sigma$ whereas cells with an angular dependence $Y_8^8(\\theta, \\phi)$ cannot have vertical velocities $> 50$ m s$^{-1}$ with a confidence level of $2\\sigma$. It is important to remember here that the GONG data sets are averaged over 108 days, while MDI data sets over 72 days. This may be somewhat larger than the lifetimes of giant cells and hence the signal may be averaged out. We can look at shorter data sets, but in that case the errors would be larger. Further, the convective cells are not expected to have smooth velocity profiles of the form used by us over the entire solar surface. That would also reduce the contribution to splitting coefficients.  Finally, we would like to point out that while the degenerate perturbation theory can be easily adopted for inversion because of linearity of effect and the form of perturbation which can be represented by appropriate kernels. In contrast, the effect of quasi-degenerate perturbation theory is nonlinear in velocity and further is mainly governed by frequency differences between nearly degenerate modes, rather than the matrix elements which depend on velocity profile. Hence it is difficult to use these shifts in inversions. For example, the main effect of flow with $s=8$ is felt in modes with turning points close to $0.7\\Rs$, which happens to be close to the lower limit of region in which our flows were localized. But we have verified that this is just a coincidence by performing calculations with different lower boundary. The same set of modes is affected irrespective of lower boundary, because it is these modes that have small frequency differences. This is only determined by the value of $s$ which determines the range of $l$ values that are coupled. Thus various types of perturbations with the same coupling will give large splittings in these modes and it is difficult to identify the actual source of perturbation."
    },
    "0910/0910.1524_arXiv.txt": {
        "abstract": "We perform hydrodynamical simulations of the accretion of pebbles and rocks onto protoplanets of a few hundred kilometres in radius, including two-way drag force coupling between particles and the protoplanetary disc gas. Particle streams interacting with the gas far out within the Hill sphere of the protoplanet spiral into a prograde circumplanetary disc. Material is accreted onto the protoplanet due to stirring by the turbulent surroundings. We speculate that the trend for prograde rotation among the largest asteroids is primordial and that protoplanets accreted 10\\%--50\\% of their mass from pebbles and rocks during the gaseous solar nebula phase. Our model also offers a possible explanation for the narrow range of spin periods observed among the largest bodies in the asteroid and trans-Neptunian belts, and predicts that 1000 km-scale Kuiper belt objects that have not experienced giant impacts should preferentially spin in the prograde direction. ",
        "introduction": "Planets form in protoplanetary gas discs as dust grains collide and grow to ever larger bodies \\citep{Safronov1969,Dominik+etal2007}. Accumulation of differentiated meteorite parent bodies must occur while $^{26}$Al and $^{60}$Fe are still present in the solar nebula, constraining their formation time to at most a million years after the formation of calcium-aluminium-rich inclusions (\\citealt{Bizzarro+etal2005}; \\citealt*{Yang+etal2007}). This coincides with the gaseous solar nebula epoch and indicates that planetesimals and protoplanets of several hundred kilometres in size were accumulated in a dense gaseous environment.  The alignment between the orbits of the planets and the rotation of the Sun is one of the key pieces of evidence for the nebular hypothesis of the formation of the solar system. The large majority of satellites also orbit their host planet in what is termed the prograde direction, i.e.\\ orbit in the same direction as the planet orbits the sun, suggesting that they formed in circumplanetary discs. This regularity seems to extend to the predominantly prograde rotations of the planets and the largest asteroids. While the gas giant planets have likely been spun up by gas accretion, growth of solid planetary bodies by accretion of km-sized planetesimals is not generally expected to lead to systematic prograde rotation. Post-accretion, giant impacts which can significantly change the obliquity and spin rate of protoplanets \\citep{DonesTremaine1993b,KokuboIda2007} do so in a stochastic way yielding no preference for prograde motion. Successful attempts to explain a {\\it systematic} prograde contribution from accretion of solids in a gas-free medium require special surface density profiles \\citep{LissauerKary1991} or ``semicollisional'' accretion \\citep{SchlichtingSari2007}, i.e. collisions between small particles and subsequent accretion into a prograde circumplanetary particle disc that feeds the central object. \\begin{figure*} \\begin{center} \\includegraphics[width=12.0cm]{fig1.eps} \\end{center} \\caption{Accretion of pebbles onto a protoplanet fixed in the centre of the coordinate system. The first panel indicates the Hill radius and the accretion radius (inner hole). The colour coding shows the local density of solids, in units of the mid-plane gas density. Particles are initially accreted from the second and fourth quadrants (second panel). As pebble streams drag the gas down towards the protoplanet, the incompressible gas expands in great particle-free bubbles, eventually making the system turbulent (third panel). A prograde accretion disc, fed from the second and fourth quadrants, forms around the protoplanet (fourth panel).} \\label{f:mainresult} \\end{figure*} \\begin{figure*} \\begin{center} \\includegraphics[width=0.7\\linewidth]{fig2.eps} \\end{center} \\caption[]{Delivery of prograde angular momentum to a protoplanet of 400 km radius (if placed at $r_0=10\\,{\\rm AU}$). The dots show the angular momentum of a subset of the accreted particles, while the blue (black) and red (grey) lines indicate the instantaneous and cumulative average. ({\\bf A}-{\\bf B}) Accretion of pebbles from environment with mean solids-to-gas ratio $\\epsilon=1.0$ and $\\epsilon=10.0$, respectively. The presence of a Keplerian particle disc is clearly visible at late times where the accreted angular momentum is exactly equal to that of a Keplerian disc (dashed line). ({\\bf C}) In the low solids-to-gas ratio case ($\\epsilon=0.1$) the pebbles have little influence on the gas and their motion is strongly damped as they approach the protoplanet, imparting only little rotation to the protoplanet. ({\\bf D}) Without drag forces these effectively very large particles give a net retrograde rotation to the protoplanet.} \\label{f:angmom_deliv} \\end{figure*}  In this paper we show that prograde rotation may be a natural consequence of planetesimal growth from accretion of cm-sized solids in a gaseous environment. Providing a physical mechanism for prograde rotation in turn supports that the dominance of prograde rotation among the largest asteroids is significant despite their low number. Our model assumes simultaneous presence of cm-sized ``pebbles'' and gas. Such conditions are supported by observations of protoplanetary discs, which show evidence of dust growth to at least a few cm \\citep{Testi+etal2003,Sicilia-Aguilar+etal2007}. The inferred population of mm-cm pebbles can be very substantial, on the order of $10^{-3}$ solar masses \\citep{Wilner+etal2005}.  We present hydrodynamical computer simulations of the accretion of pebbles and rocks onto young protoplanets in a gaseous nebula environment. We find that particle streams interacting with the gas far out within the Hill sphere dissipate enough energy that the particles spiral down towards the protoplanet and form a prograde accretion disc. The prograde motion is transferred to protoplanet rotation as the particles are accreted. To explain the observed rotations of the largest asteroids and Kuiper belt objects, we propose that protoplanets accreted a major fraction (10\\%--50\\%) of their mass from pebbles and rocks during the gaseous solar nebula phase.  The paper is organised as follows. In \\S\\ref{s:setup} we describe the set up of our simulations. In \\S\\ref{s:results} the main results are presented, both for 2-D and 3-D simulations. We speculate in \\S\\ref{s:discussion} about implications for the observed prograde rotation of large asteroids and the future prospect of measuring rotation directions of Kuiper belt objects. We conclude on our findings in \\S\\ref{s:conclusions}. The appendices contain details that do not appear in the main paper and several tests of the validity of our results. In Appendix \\ref{s:dyneq} the dynamical equations are presented. We validate the code in Appendix \\ref{s:validate} against known results on the accreted angular momentum. Appendix \\ref{s:mmsn} describes drag forces in more detail and also the translation from dimensionless friction time to physical particle size. In Appendix \\ref{s:collide} we compare time-scales of drag force and particle collisions and conclude that drag forces are by far the dominant force, except in regions that are highly particle dominated. In Appendix \\ref{s:converge} we present results at lower and higher resolution and conclude that prograde rotation is strengthened with increasing resolution. The dependence of the results on the size of the simulation box are treated in Appendix \\ref{s:boxsize}. ",
        "conclusions": "\\label{s:conclusions}  We present computer simulations of the accretion of pebbles and rocks onto young protoplanets in a gaseous nebula environment. Our findings can be summarised as follows:  \\begin{itemize} \\item Particle streams interacting with the gas far out in the Hill sphere of the growing protoplanet dissipate enough energy that the particles spiral down towards the central object and form a prograde accretion disc. The prograde motion is transferred to the protoplanet as the particles are accreted. \\item In order to explain the observed rotations of the largest asteroids and Kuiper belt objects, we propose that protoplanets accreted a major fraction (10\\%--50\\%) of their mass from pebbles and rocks during the gaseous solar nebula phase. \\item A specific prediction of our model is that most of the largest Kuiper belt objects should possess prograde spins. \\end{itemize}  Since the Keplerian surface frequency depends only on the density of the protoplanet, accretion from a Keplerian particle disc explains how asteroids and Kuiper belt objects can have similar rotation periods even though the orbital time-scale is much longer in the Kuiper belt. Our results imply not just a universality in the rotation period of solid bodies in the solar system, but also in the formation process and in the important role of gas and pebbles for the growth of hundred kilometre scale protoplanets.  Future numerical work should focus on the extension of the presented model to include turbulence driven by global instabilities and more realistic solid physics, such as a distribution of particle sizes, collisional viscosity and collisional fragmentation."
    },
    "0910/0910.1234.txt": {
        "abstract": "{We investigate the regularity of cluster pressure profiles with  \\rexcess, a representative sample of 33 local ($z < 0.2$) clusters drawn from the \\reflex\\ catalogue and observed with \\xmm. The sample spans a mass range of $10^{14}\\msol<\\Mv<10^{15}\\msol$, where $\\Mv$ is the mass corresponding to a density contrast of 500. We derive an average profile from observations scaled by mass and redshift according to the standard self-similar model, and find that the dispersion about the mean is remarkably low, at less than 30 per cent beyond $0.2\\,\\Rv$, but increases towards the centre. Deviations about the mean are related to both the mass and the thermo-dynamical state of the cluster.  Morphologically disturbed systems have systematically shallower profiles while cooling core systems are more concentrated. The scaled profiles exhibit a residual mass dependence with a slope of $\\sim 0.12$, consistent with that expected from the empirically-derived slope of the \\MY\\ relation; however, the departure from standard scaling decreases with radius and is consistent with zero at $R_{500}$. The scatter in the core and departure from self-similar mass scaling is smaller compared to that of the entropy profiles, showing that the pressure is the quantity least affected by dynamical history and  non-gravitational physics. Comparison with scaled data from several state of the art numerical simulations shows good agreement outside the core. Combining the observational data in the radial range $[0.03$--$1]\\,\\Rv$ with simulation data in the radial range $[1$--$4]\\,\\Rv$, we derive a robust measure of the universal pressure profile, that, in an analytical form, defines the physical pressure profile of clusters as a function of mass and redshift up to the cluster 'boundary'. Using this profile and direct spherical integration of the observed pressure profiles,  we estimate the integrated Compton parameter $Y$ and investigate its scaling with $\\Mv$ and $\\LX$, the soft band X--ray luminosity. We consider both the spherically integrated quantity, $Y_{\\rm sph}(R)$, proportional to the gas thermal energy, and the cylindrically integrated quantity, $Y_{\\rm cyl}(R)=\\YSZ D_{\\rm A}^2$, which is directly related to the Sunyaev-Zel'dovich (SZ) effect signal. From the low scatter of the observed $Y_{\\rm sph}(\\Rv)$--$\\YX$ relation we show that  variations in pressure profile shape do not introduce extra scatter into the  $Y_{\\rm sph}(\\Rv)$--$\\Mv$ relation as compared to that from the $\\YX$--$\\Mv$ relation.  The $Y_{\\rm sph}(\\Rv)$--$\\Mv$ and $Y_{\\rm sph}(\\Rv)$--$\\LX$ relations derived from the data are in excellent agreement with those expected from the universal profile. This profile is used to derive the expected $\\YSZ$--$\\Mv$  and \\YLX\\  relations for any aperture. } ",
        "introduction": "Galaxy clusters provide valuable information on cosmology, from the nature of dark energy to the physics driving galaxy and structure formation. Clusters are filled with a hot ionised gas that can be studied both in X-ray and through the thermal Sunyaev-Zel'dovich  (SZ) effect, a spectral distortion of the cosmic microwave background (CMB) generated via inverse Compton scattering of CMB photons by the free electrons. Its magnitude is proportional to the Compton parameter $y$, a measure of  the gas pressure integrated along the line-of-sight, $y = (\\sigma_{\\rm T}/m_{\\rm e} c^2)\\int P dl$, where $\\sigma_{\\rm T}$ is the Thomson cross-section, $c$ the speed of light,  $m_{\\rm e}$ the electron rest mass and $P=\\ne T$ is the product of the electron number density and temperature. The total SZ signal, integrated over the cluster extent, is proportional to the integrated Compton parameter $\\YSZ$, $\\YSZ D_{\\rm A}^2 = (\\sigma_{\\rm T}/m_{\\rm e} c^2)\\int P dV$, where $D_{\\rm A}$ is the angular distance to the system.  As the gas pressure is  directly related to the depth of the gravitational potential,  $\\YSZ D_{\\rm A}^2$  is expected to be closely related to the mass. Numerical simulations \\citep[e.g.,][]{das04, mot05, nag06, bon07} and analytical models \\citep{rei06} of cluster formation indicate that the intrinsic scatter of the \\YM\\ relation is low, regardless of the cluster dynamical state \\citep[see also][]{wik08} or the exact details of the gas physics. However, the normalisation of the relation {\\it does} depend on the gas physics \\citep{nag06,bon07}, as does the exact amount of scatter, the details of which are still under debate  \\citep{sha08}. Given that this relation, and the underlying pressure profile, are key ingredients for the use of on-going or future SZ cluster surveys for cosmology, and provide invaluable information on the physics of the intra-cluster medium (ICM), it is important to calibrate these quantities from observations.  In recent years, SZ observational capability has made spectacular progress, from the first spatially resolved  (single--dish) observations of individual objects \\citep[][]{poi99,poi01,kom99, kom01} to the first discovery of new clusters with a blind SZ survey \\citep[][]{sta09}. Spatially resolved SZE observations directly probe  the mass weighted temperature along the line of sight. By contrast,  temperatures derived from X-ray spectra, by fitting an isothermal model to a multi-temperature plasma emission along the line of sight, are likely to be biased \\citep{mat01}. Although schemes to correct for this effect have been defined  \\citep{maz04,vik06b}, it remains a potential source  of systematics.  Stacking analysis of WMAP data  around known X--ray clusters has allowed statistical detection of a scaled pressure profile \\citep{afs07} or a spatially resolved decrement \\citep{lie06,atr08,die09}, showing clear discrepancies with the prediction of a simple isothermal $\\beta$--model.  Pressure or temperature profiles of individual clusters have started to be derived from combined analysis of X-ray and SZE imaging data, using non-parametric deprojection methods \\citep{nor09} or more realistic  models than the $\\beta$-model \\citep{kit04, mro09}.   Interestingly, the profiles are found to be consistent with profiles derived using X-ray spectroscopic data \\citep[see also][]{jia08,hal09}. However, such studies are still restricted to a few test cases,  particularly  hot  clusters.  The \\YM\\ relation has been recently derived by \\citet{bon08},  an important step forward as compared to previous work  based  on central decrement measurements using heterogenous data sets \\citep{mcc03, mor07}; however, quantities were estimated within  $R_{2500}\\sim 0.44\\Rv$ \\footnote{Here and in the following, $M_{\\delta}$ and $R_{\\delta}$ are the total mass and radius corresponding to a density contrast, $\\delta$, as compared to $\\rho_{\\rm c}(z)$,   the critical density of the universe at the cluster redshift: $M_{\\delta}= (4\\pi/3) \\delta  \\rho_c(z)  R_{\\delta}^3$. $\\Mv$ corresponds roughly to the virialised portion of clusters, and is traditionally used to define the 'total' mass.} and assuming an isothermal $\\beta$--model, which may provide a biased estimate \\citep{hal07}. In addition, the first scaling relation using weak lensing masses, rather than X--ray hydrostatic masses, has now appeared \\citep{mar09}, although constraints from these data are currently weak. \\\\  In this context, statistically more precise, albeit indirect, information can be obtained from X-ray observations.  A key physical parameter is $\\YX$,  the X--ray analogue  of the integrated Compton parameter,  introduced by \\citet{kra06}.  $\\YX$ is defined as the product  of  $\\Mgv$,  the gas mass within $\\Rv$ and $\\TX$, the spectroscopic temperature outside the core. The local \\MY\\ relation for relaxed clusters has recently been calibrated \\citep{nag07,mau07,app07, vik09}, with excellent agreement achieved between various observations  \\citep[e.g., see][]{app07}.  However, the link between $\\YX$ and $\\YSZ$ depends on cluster structure through % \\begin{equation} \\frac{\\YSZ D_{\\rm A}^2 }{\\YX} =  \\frac{\\sigma_{\\rm T}}{m_{\\rm e} c^2}\\frac{1}{ \\mu_{\\rm e} m_{\\rm p}} \\frac{\\langle \\ne T\\rangle} {\\langle\\ne\\rangle_{\\Rv} \\TX} \\label{eq:yszyx} \\end{equation} % \\noindent where the angle brackets denote volume averaged quantities. From Eq.~\\ref{eq:yszyx}, it is clear that an understanding of the radial pressure distribution and its scaling is important not only as a probe of the ICM physics, but also for exploitation of these data. High resolution measurements of the radial density and temperature distribution are now routinely available from X--ray observations  but the pressure profile structure and scaling have been relatively little studied. The pressure profiles of groups have been studied by \\citet{fin06} and \\citet{joh09}. In the cluster regime,  \\citet{fin05} analysed the 2D  pressure distribution in a flux-limited sample of 6 hot ($\\kT>7 \\keV$) clusters at $z\\sim0.3$ showing  fluctuations at the $30\\%$ level around the mean profile, scaled by temperature.  To our knowledge, the only study of pressure profiles scaled by mass is that of \\citet{nag07}, who used {\\it Chandra} X-ray observations to derive a universal pressure profile, with the external slope derived from numerical simulations. However, their sample was restricted to hot ($kT>5\\keV$) relaxed clusters, which are all  cool  core systems, and contained five objects. For the reasons mentioned above, it is of considerable interest to extend this analysis to data from a larger and more representative sample of the cluster population.  In this paper we do this by investigating the regularity of cluster pressure profiles with  \\rexcess\\  \\citep{boh07}, a representative sample of 33 local ($z < 0.2$) clusters drawn from the \\reflex\\ catalogue \\citep{boh04} and observed with \\xmm. We derive an average profile from observations scaled by mass and redshift according to the self-similar model and  relate the deviations about the mean to both the mass and the thermo-dynamical state of the cluster (Sec.~\\ref{sec:psc}). Comparison with data from several state of the art numerical simulations  (Sec.~\\ref{sec:simul}) shows good agreement outside the central regions, which is the most relevant aspect for the $\\YSZ$ estimate. Combining the observational data in the radial range $[0.3$--$1]\\,\\Rv$ with simulation data in the radial range $[1$--$4]\\,\\Rv$  allows us to derive a robust measure of the universal pressure profile up to the cluster 'boundary' (Sec.~\\ref{sec:puniv}). Using this profile or direct spherical integration of the observed pressure profiles,  we estimate the spherically and cylindrically integrated Compton parameter and investigate its scaling with $\\YX$, $\\Mv$ and $\\LX$, the soft band X--ray luminosity (Sec.\\ref{sec:ysz}).  We adopt a $\\Lambda$CDM cosmology with $H_{0} = 70$~ km/s/Mpc, $\\Omega_{\\rm M} = 0.3$ and $\\Omega_\\Lambda= 0.7$. $h(z)$ is the ratio of the Hubble constant at redshift $z$ to its present value, $H_{0}$. $\\TX$ is the temperature  measured in the $[0.15$--$0.75]~\\Rv$ aperture.  All scaling relations are derived using the BCES orthogonal regression method with bootstrap resampling \\citep{akr96}, and uncertainties are quoted throughout at the 68 per cent confidence level.  %----------------------------------------------------------------------------------------------------------------------------------------------------------------------------------- %----------------------------------------------------------------------------------------------------------------------------------------------------------------------------------- ",
        "conclusions": "The present work is the first examination of the properties of the ICM pressure for a representative sample of nearby clusters. The sample, \\rexcess, was chosen by X-ray luminosity alone, without regard to morphology or dynamical state. It  covers the mass range $10^{14} < M_{500} < 10^{15}$ M$_\\odot$, with mass iteratively estimated from the \\MY\\ relation, calibrated from a sample of relaxed clusters including \\rexcess\\ objects.  As for the entropy \\citep{pra09b}, the depth of the observations allowed us to probe the scaling behavior of the pressure profiles out to $\\Rv$. This is essential for a complete picture of the modification of the standard self-similarity due to non-gravitational processes, including its radial behavior. \\\\  Scaling the individual pressure profiles by mass and redshift according to the standard self-similar model, we derived an average scaled pressure profile for the cluster population and  relate the deviations about the mean to both the mass and the thermo-dynamical state of the cluster: \\begin{itemize} \\item Cool core systems exhibit more peaked profiles, while morphologically disturbed systems have shallower profiles. \\item As a result, the dispersion is large in the core region, reaching approximately 80 per cent at $0.03\\,\\Rv$.  However, as compared to the density, the pressure exhibits less scatter, a result of the anticorrelation of the density and temperature profiles interior to $0.2\\,\\Rv$. Outside the core regions, the dispersion about the average profile is remarkably low, at less than 30 per cent beyond $0.2\\,\\Rv$. \\item We find  a residual mass dependence of the scaled profiles, with a slope of $\\sim 0.12$, consistent with that expected from the empirical non-standard slope of the \\MY\\ relation. However, there is some evidence that  the departure from standard scaling decreases with radius and is consistent with zero at $R_{500}$.  We provide an analytical correction to the mean slope that accounts for this second order effect. \\end{itemize} The behaviour of the pressure profiles,  with respect to  standard self-similarity with zero dispersion, resembles that generally found for other quantities such as the entropy or density: 1) regularity in shape outside the core 2) increased dispersion inside the core linked to cooling effects and dynamical state and 3) departure from standard mass scaling that becomes less pronounced towards the cluster outskirts. However, the latter two  deviations are less pronounced than for the entropy and/or density, showing that the pressure is the quantity least affected by dynamical history and  non-gravitational physics. This further supports the view that  $Y_{\\rm SZ}$ is indeed a  good mass-proxy.  Furthermore, our direct measure of the  $Y_{\\rm sph}(\\Rv)$--$\\YX$ relation, where  $Y_{\\rm sph}(\\Rv)$ is the spherically integrated pressure profile,  exhibits  dispersion consistent with the $<5\\%$ statistical scatter.  This shows that  variations in pressure profile shape do not introduce significant extra intrinsic scatter into the  $Y_{\\rm sph}(\\Rv)$--$\\Mv$ relation as compared to that from the $\\YX$--$\\Mv$ relation. \\\\  The observational data are compared to and combined  with simulated data to derive the universal ICM pressure profile. This profile is then used to predict  the scaling relations involving the integrated Compton parameter $Y$. We consider both the spherically integrated quantity, $Y_{\\rm sph}(R)$, which is related to the gas thermal energy, and also the cylindrically integrated quantity $Y_{\\rm cyl}(R) = Y_{SZ}(R) D_A^2$, which is directly related to the observationally-derived SZ signal within $\\theta=R/D_{\\rm A}$: \\begin{itemize} \\item Simulated scaled profiles from three independent sets of state of the art numerical simulations show excellent agreement, within $20\\%$, between $0.1$ and $3\\Rv$, for pressures varying by 4 orders of magnitude in that radial range. \\item Comparison with observed scaled data shows good agreement outside the core regions, which is the most relevant aspect for the $Y_{\\rm SZ}$ estimate. The average simulation profile lies parallel to the observed data, with only a slight offset ($\\sim 10$ per cent)  when the simulated profiles are scaled using the hydrostatic mass. \\item This motivates us to combine the average observed  scaled profile in the $[0.03-1]\\,\\Rv$ radial range with the average simulated profile  in the $[1-4]\\,\\Rv$ range. This hybrid profile is fitted by a generalised NFW model, which allows us to define a dimensionless universal ICM pressure profile. Combined with  the empirical mass scaling of the profiles, this universal profile defines the physical pressure profile of clusters, up to the cluster boundary, as a function of mass and redshift, assuming self-similar evolution. \\item  This universal profile allows us to derive the expected $Y_{\\rm sph}(x\\Rv)$--$\\Mv$  or $\\YSZ(x\\Rv)$--$\\Mv$  relations for any aperture. The slope is the inverse of the empirical slope of the \\MY\\ relation. The normalisation is given by the dimensionless integral of the universal profile within the region of interest expressed in scaled radius.  The corresponding  \\YLX\\  relations can be derived by combining the relevant $Y$--$\\Mv$ relation with the empirical $\\LX$--$\\Mv$ relation. \\item  The $Y_{\\rm sph} (\\Rv)$--$ \\Mv$ and $Y_{\\rm sph} (R_{500})$-- $\\LX$ relations derived directly from the individual profiles are in excellent  agreement with those expected from the universal profile. \\item We confirm that the isothermal $\\beta$--model over-estimates the $Y$ signal at given mass. This overestimate depends strongly on the assumption on cluster extent and reaches a factor of nearly two at $2\\Rv$. \\end{itemize} The convergence  of  various approaches to determine scaled  cluster profiles supports the robustness of our determination of the  universal pressure profile, particularly its shape. This includes the agreement  between independent simulations, between these simulations and the present observed data based on a representative cluster sample, and also the agreement between the  present \\xmm\\ data and published  \\chandra\\ data for clusters of similar thermo-dynamical state. As a result, we believe that  quantities which purely depend on the universal profile shape are particularly robust and well converged. This includes the typical SZ decrement  profile or relations between the Compton parameter  estimated in various apertures.  However, the pressure profile interior to  $\\Rv$ is derived from temperatures estimated using azimuthally averaged spectra. These have been corrected for the spectroscopic bias due to projection but not for azimuthal variations. In the cluster outskirts,  the electron-proton equilibration time is larger than the Hubble time \\citep{fox97} and if the electron temperature is indeed smaller than the ion temperature, this will affect the pressure profile and  lead to a decrease in the total $Y_{SZ}$ signal \\citep{rud09}. High resolution SZ data  with improved sensitivity are needed to probe any remaining systematic effects due to the spectroscopic bias, and to directly observe the shape of the pressure profile beyond $\\Rv$, which is out of reach of current X--ray observatories.  Using the universal profile, the absolute normalisation and slope of the $Y$--$\\Mv$ relations rely on  the underlying  observationally defined  \\MY\\ relation.  Initial comparison with $Y_{\\rm SZ}(\\Rv)$ data for 3 high mass systems,  measured with SZA by \\citet{mro09} and  analysed with a realistic analytic pressure profile, indicates good agreement.  A key point is to extend this type of analysis to larger samples and  to include lower mass systems.  We further emphasize that the  \\MY\\ relation  was calibrated from  hydrostatic mass estimates  using relaxed objects. The $Y$--$\\Mv$ relation we derive is  technically a $Y$--X--ray mass relation and  is expected to differ from the `true' $Y$--$\\Mv$ by the offset between the `true' mass and the  hydrostatic mass for {\\it relaxed} objects.  A major open issue is the pressure evolution.  With the present study based on a local cluster sample, we could only assume standard self-similar evolution.  Because the SZ signal is not subject to redshift dimming, on going SZ surveys are expected to detect many new clusters at high z. Of particular interest is the  \\planck\\ survey, which, thanks to its All--Sky coverage, will detect massive, thus rare, clusters, the best objects for precise cosmology with clusters.  SZ follow-up, at the best possible resolution, and sensitive X--ray follow-up (particularly with \\xmm) will be crucial to assess  possible evolution of pressure profile shape and  measure the evolution of the \\MY\\ and $Y_{SZ}$--$\\Mv$ relations. \\\\  As a matter of practical application,  the universal pressure profile is given in Eq.~\\ref{eq:pgnfw} with parameters in Eq.~\\ref{eq:pargnfw}. For clusters of given mass $\\Mv$ and $z$, the physical pressure profile can then be derived from Eq.~\\ref{eq:puniv} and the spherical $Y_{\\rm sph} (R)$ or  cylindrical $Y_{\\rm cyl} (R)$ quantities can be estimated for any radius of interest using Eq.~\\ref{eq:mysph}--\\ref{eq:Ix} and Eq.~\\ref{eq:mysz}--\\ref{eq:Jx}, respectively.  These equations can be used as is when  $\\Mv$ is estimated for relaxed systems using the HSE equation,  and for all clusters using $\\Mv$ derived from mass-proxy  relations. The preferential relations would be the  \\MY\\ and  the  $\\Mv$--$\\LX$, where $\\LX$ is the core--excised bolometric luminosity \\citep{pra09}, as both these relations display low scatter, compared to the relation between $\\Mv$ and the full aperture soft band $\\LX$.  A typical application would be to predict the SZ signal of a known X--ray cluster with measured $\\LX$ or $\\Mv$, or to estimate the mass and thus X-ray properties of newly discovered SZ clusters. Other applications include the analysis of low S/N and/or poor resolution  SZ observation of X-ray clusters, e.g., allowing to optimise  the integration aperture and use a realistic decrement shape.  On the other hand, care is needed when knowledge of the 'true' mass is important, e.g., in predicting cluster number counts for future SZ surveys or  in SZ selection function modelling. The above total $Y_{\\rm SZ}$--$\\Mv$ relation should be corrected by the bias between the true mass and the HSE mass at $\\Rv$, which is typically $\\sim 13\\%$ as determined from comparison with current numerical simulations.  Further progress on this fundamental question, as well as on the intrinsic scatter of the $Y$--$M$ relation,  is expected from the wealth of high quality multi-wavelength data that will be available in the coming years."
    },
    "0910/0910.3717_arXiv.txt": {
        "abstract": "Several processes can cause the shape of an extrasolar giant planet's shadow, as viewed in transit, to depart from circular.  In addition to rotational effects, cloud formation, non-homogenous haze production and movement, and dynamical effects (winds) could also be important.  When such a planet transits its host star as seen from Earth, the asphericity will introduce a deviation in the transit lightcurve relative to the transit of a perfectly spherical (or perfectly oblate) planet.  We develop a theoretical framework to interpret planet shapes.  We then generate predictions for transiting planet shapes based on a published theoretical dynamical model of HD189733b.  Using these shape models we show that planet shapes are unlikely to introduce detectable lightcurve deviations (those $>~1\\times10^{-5}$ of the host star), but that the shapes may lead to astrophysical sources of systematic error when measuring planetary oblateness, transit time, and impact parameter. ",
        "introduction": "Transits are proving to be the key to characterizing extrasolar giant planets. Transit lightcurve photometry has allowed measurements of planets' orbital inclination and radius, which when combined with radial velocity observations allows an unambiguous determination of a planet's mass and density.  To date such measurements have been made for 55 transiting planets (for an up-to-date list of known transiting planets see the Extrasolar Planets Encyclopedia at {\\tt http://exoplanet.eu/catalog-transit.php/}).  Previous groups have suggested that visible-light transit photometry could determine a planet's oblateness \\citep{Seager.oblateness,oblateness.2003}, the existence of ring systems \\citep{2004ApJ...616.1193B}, and/or the potentially artificial nature of the transiting object \\citep{2005ApJ...627..534A}.  However, the published lightcurves of each of the transiting planets discovered thus far indicate no deviations from sphericity within errors.  Infrared observations of transiting planets' secondary eclipse has revealed the dayside brightness temperature of TRES-1 \\citep{2005ApJ...626..523C}, HD209458b \\citep{2005Natur.434..740D,2008ApJ...673..526K} and HD189733b \\citep{2006ApJ...644..560D,2008ApJ...686.1341C,2008Natur.456..767G} and the spectrum of HD209458b \\citep{2007Natur.445..892R,2008ApJ...674..482S} and HD189733b \\citep{2007ApJ...658L.115G}.  Numerical simulations of atmospheric dynamics predict that the photospheres of hot Jupiters are non-homogenous \\citep{2003ApJ...587L.117C,2005ApJ...629L..45C,2007ApJ...657L.113L,2008ApJ...673..513D,2008arXiv0809.2089S}. The longitudinal structure of these thermal variations was predicted to be detectable using infrared photometry over the course of a full orbit \\citep{2006ApJ...652..746F} and may also be detectable during the planet's ingress and egress from secondary eclipse \\citep{2006ApJ...649.1020W,2007ApJ...664.1199R}.  Recent \\emph{Spitzer} Space Telescope measurements have shown phase-dependent infrared flux variability for $\\mathrm{\\upsilon}$ Andromedae \\citep{2006Sci...314..623H} and HD189733 \\citep{2007Natur.447..183K,2008arXiv0802.1705K}.  These measurements provide the first constraints for dynamical models; however, the measured phase function does not match that predicted theoretically.  In this paper, we show that atmospheric dynamics also introduces asphericity to a transiting planet's observed shape, but that the lightcurve deviations thus produced are likely too small to be detected.  First we investigate the processes that might affect a planet's shape (Section \\ref{section:shape}).  Next we derive an analytical expression for a planet's shape in the simplified case of modification by a parameterized eastward equatorial jet (Section \\ref{section:analytical}).  We then develop a theoretical framework for the effects of planet shape on transit lightcurves numerically using hypothetical regularly-shaped planets as a guide (Section \\ref{section:numerical}).  Finally we calculate the predicted shape of HD189733b based on the new dynamical models of \\citet{2008arXiv0809.2089S} (Section \\ref{section:practical}), and calculate the transit lightcurves for the predicted planet shape models, discussing shape detectability from ground- and space-based photometry. ",
        "conclusions": "\\label{section:conclusion}  Planetary winds affect the three-dimensional shape of a planet's  constant-density contours, leading to departures from sphericity.  A planet's silhouette as viewed in transit should depend on the projected shape of these constant-density surfaces, assuming that absorption depends only on atmospheric density.  The resulting silhouette should affect the planet's transit lightcurve.  We calculate an analytical estimate for the amplitude of the shape variation, using simplifying assumptions regarding a planet's zonal wind structure.  Using the maximum equatorial winds found in dynamical models of HD209458b \\citep{2005ApJ...629L..45C} predicts a atmospheric-dynamics-driven global equator-to-pole radius difference of between 480 km and 960~km.  Using results from the more sophisticated \\citet{2008arXiv0809.2089S} model of HD189733b, the analytic expression predicts a 715~km atmospheric-dynamics-driven global equator-to-pole radius difference.  For a more robust estimate of a planet's transit shape, we use the \\citet{2008arXiv0809.2089S} model directly by deriving the wind-induced shape from contours of constant density along the planet's terminator.  The resulting shape has a total radius difference of around $\\sim700$~km.  However, when we use our numerical lightcurve-fitting routine to estimate the detectability of this shape in a transit lightcurve, we find that it is not detectable, with lightcurve residuals of order only $10^{-5}$ of the stellar flux.  To understand why the detectability is so low, we calculate detectabilities for planets with regular shapes such that $r(\\theta)=R_p+A\\sin(n\\Theta+\\Psi)$.  The $n=0$ term corresponds to the planet's average radius, the $n=1$ term to offsets between the center of mass and center of figure, the $n=2$ term to planetary oblateness, and higher order terms to more complex shapes.  The $n=1$ term is not detectable in a transit lightcurve. The projected shapes that are symmetric with respect to the planet's orbit normal produce symmetric lightcurves; those shapes that are asymmetric produce antisymmetric lightcurves. In general shapes with lower $n$ have higher detectabilities, and require less-fine time resolution than shapes with higher $n$.  However the symmetric $n=2$ term has somewhat low detectability due to the ability of a spherical planet model to partially mimic its transit lightcurve signature.  A Fourier decomposition of the model HD189733b shape reveals why it would be so hard to detect.  Most of the amplitude of the shape is in the $n=1$ and $n=2$ terms:  the $n=1$ is undetectable, and the $n=2$ has relatively low detectability.  The higher order terms become progressively harder to detect as $n$ increases.  Thus the smooth nature of the predicted HD189733b shape leads to its low transit detectability.  The detectability of other planets will depend on their wind velocities and zonal wind structures.  However, the high insolation and low rotation rates of all hot Jupiters may drive them to behave similarly to HD189733b.  Multiple counterrotating jets would avoid the low detectability of the oblateness ($n=2$) term, but higher $n$ terms have lower detectabilities as well.  Variation in the height of clouds or the absorption from haze or the atmosphere around the disk could lead to detectable shapes, but only in unusual circumstances.  Hence we think that it is unlikely that the winds on transiting planets will be able to affect their lightcurves at a level that will be detectable in the near-future."
    },
    "0910/0910.0839_arXiv.txt": {
        "abstract": "I describe the conceptual and mathematical basis of an approach which describes gravity as an emergent phenomenon. Combining the principle of equivalence and the principle of general covariance with known properties of local Rindler horizons, perceived by observers accelerated with respect to local inertial frames, one can provide a thermodynamic re-interpretation of the field equations describing gravity in any diffeomorphism invariant theory. This fact, in turn, leads us to the possibility of deriving the field equations of gravity by maximising a suitably defined entropy functional, without using the metric tensor as a dynamical variable. The approach synthesizes concepts from quantum theory, thermodynamics and gravity leading to a fresh perspective on the nature of gravity. The description is presented here in the form of a dialogue, thereby addressing several frequently-asked-questions. ",
        "introduction": "\\ha \\footnote{Harold was  a very useful creation originally due to Julian Schwinger \\cite{julian} and stands for Hypothetically Alert Reader Of Limitless Dedication. In the present context, I think of Harold as Hypothetically Alert Relativist  Open to Logical Discussions.} For quite sometime now, you have been talking about `gravity being an emergent phenomenon' and a `thermodynamic perspective on gravity'. This is  quite different from the conventional point of view in which gravity is a fundamental interaction and spacetime thermodynamics of, say, black holes is a particular result which can be derived in a specific context. Honestly, while I find your papers fascinating I am not clear about the broad picture you are trying to convey. Maybe you could begin by clarifying what this is all about, before we plunge into the details ? What is the roadmap, so to speak ?  \\me To begin with, I will show you that the equations motion describing gravity in  \\textit{any diffeomorphism invariant theory} can be given \\cite{tp09papers} a  suggestive thermodynamic re-interpretation (Sections \\ref{sec:lro}, \\ref{sec:reinterpret}). Second, taking a cue from this, I can formulate a variational principle for a suitably defined entropy functional --- involving both gravity and matter --- which will lead to the field equations of gravity \\cite{aseementropy,tpbojo} without varying the metric tensor as a dynamical variable (Section \\ref{sec:eqnnewvar}).  \\ha Suppose I have an action for gravity plus matter (in $D$ dimensions) \\begin{equation} A=\\int d^Dx \\sqrt{-g}\\left[L(R^{ab}_{cd}, g^{ab})+L_{matt}(g^{ab},q_A)\\right] \\end{equation} where $ L$ is any scalar built from metric and curvature  and $L_{matt}$ is the matter Lagrangian depending on the metric and some matter variables $q_A$. (I will assume $L$ does not involve derivatives of curvature tensor, to simplify the discussion.) If I vary $g^{ab}$ in the action I will get some  equations of motion (see e.g. Refs. \\cite{ayan,mohut}), say,  $2E_{ab}=T_{ab}$ where $E_{ab}$ is \\footnote{The signature is - + + + and Latin letters cover spacetime indices while Greek letters run over space indices.} \\begin{equation} E_{ab}=P_a^{\\phantom{a} cde} R_{bcde} - 2 \\nabla^c \\nabla^d P_{acdb} - \\frac{1}{2} L g_{ab};\\quad P^{abcd} \\equiv \\frac{\\partial L}{\\partial R_{abcd}} \\label{genEab} \\end{equation} Now, you are telling me that (i) you can give a thermodynamic interpretation to the equation $2E_{ab}=T_{ab}$ just because it comes from a scalar Lagrangian and (ii) you can also derive it from an entropy maximisation principle. I admit it is  fascinating. But why should I take this  approach as more fundamental, conceptually, than the good old way of just varying the total Lagrangian $L+L_{matt}$ and getting $2E_{ab}=T_{ab}$ ? Why is it more than a curiosity ?  \\me That brings me to the third aspect of the formulation which I will discuss towards the end (Section \\ref{sec:comaparison}). In my approach, I can provide a natural explanation to several puzzling aspects of gravity and horizon thermodynamics all of which have to be thought of as mere algebraic accidents in the conventional approach you mentioned. Let me give an analogy. In Newtonian gravity, the fact that inertial mass is equal to the gravitational mass is an algebraic accident without any fundamental explanation. But in a geometrical theory of gravity based on principle of equivalence, this fact finds a natural explanation. Similarly, I think we can make progress by identifying key facts  which have no explanation in the conventional approach and providing them  a natural explanation from a different perspective. You will also see that this approach connects up several pieces of conventional theory  in an elegant manner.   \\ha Your ideas also seem to be quite different from other works which describe gravity as an emergent phenomena \\cite{analog}. Can you explain \\textit{your} motivation?   \\me Yes. The original inspiration for my work, as for many others, comes from the old idea of Sakharov \\cite{sakharov} which attempted  to describe spacetime dynamics as akin to the theory of elasticity. There are two crucial differences between my approach and many other ones.  To begin with, I realized that the thermodynamic description transcends Einstein's general relativity and can incorporate a much wider class of theories --- this was first pointed out in ref. \\cite{paris} and elaborated in several of my papers --- while many other approaches concentrated on just Einstein's theory. In fact, many other approaches use techniques strongly linked to Einstein's theory -- like for example, Raychaudhuri equation to study rate of change of horizon area, which is difficult to  generalize to theories in which the horizon entropy is \\textit{not} proportional to horizon area. I use more general techniques.  Second, I work at the level of action principle and its symmetries to a large extent  so I have a handle on the off-shell structure of the theory; in fact, much of the thermodynamic interpretation in my approach is closely linked to the structure of action functional (like e.g., the existence of surface term in action, holographic nature etc.) for gravitational theories. This link is central to me while it is not taken into account in any other approach.   \\ha So essentially you are claiming that the thermodynamics of horizons is more central than the dynamics of the gravitational field while the conventional view is probably the other way around. Why do you stress the thermal aspects of horizons so much ? Can you give a motivation ?  \\me Because thermal phenomenon is a window to microstructure! Let me explain. We know that the continuum description of  a fluid, say,  in terms of a set of dynamical variables like density $\\rho$, velocity $\\mathbf {v}$,  etc. has a  life of its own. At the same time, we also know that these dynamical variables and the description have no validity at a fundamental level where the matter is discrete. But one can actually \\textit{guess} the existence of microstructure without using any experimental proof for the  molecular nature of the fluid, just from the fact that the fluid or a gas  exhibits \\textit{thermal phenomena} involving temperature and transfer of heat energy. If the fluid is treated as a continuum and is described by  $\\rho (t, \\mathbf {x})$,  $ \\mathbf {v}(t, \\mathbf {x})$ etc., all the way down,   then it is \\textit{not} possible to explain the thermal phenomena in a natural manner. As first stressed by Boltzmann, the heat content of a fluid arises due to random motion of discrete microscopic structures which \\textit{must} exist in the fluid. These new degrees of freedom --- which we now know are related to the actual molecules --- make the fluid capable of storing energy internally and exchanging it with surroundings. So, given an apparently continuum phenomenon which exhibits temperature, Boltzmann could infer the existence of underlying discrete degrees of freedom.  \\ha I agree. But what does it lead to in the present context?  \\me The paradigm is: \\textit{If you can heat it, it has microstructure!} And you can heat up spacetimes by collapsing matter or even by just accelerating \\cite{daviesunruh}. The horizons which arise in general relativity are endowed with temperatures \\cite{hawking} which shows that, at least in this context, some microscopic degrees of freedom are coming into play. So  a thermodynamic description that links the standard description of gravity with the statistical mechanics of --- as yet unknown --- microscopic degrees of freedom must exist. It is in this sense that I consider gravity to be emergent.   Boltzmann's insight about the thermal behaviour  has two other attractive features which are useful in our context. First, while the existence of the discrete degrees of freedom is vital in such an approach, the exact nature of the degrees of freedom is largely irrelevant. For example, whether we are dealing with argon molecules or helium molecules is largely irrelevant in the formulation of  gas laws and such differences can be taken care of in terms of a few well-chosen numbers (like, e.g.,  the specific heat). This suggests that such a description will have certain amount of robustness and independence as regards the precise nature of  microscopic degrees of freedom.  Second, the entropy of the system arises due to our ignoring the microscopic degrees of freedom. Turning this around, one can expect the form of entropy functional to encode the essential aspects of microscopic degrees of freedom, even if we do not know what they are. If we can arrive at the appropriate form of entropy functional, in terms of some effective degrees of freedom, then we can expect it to provide the correct description.\\footnote{Incidentally, this is why thermodynamics needed no modification due to either relativity or quantum theory. An equation like $TdS=dE+PdV$ will have universal applicability as long as effects of relativity or quantum theory are incorporated in the definition of $S(E,V)$ appropriately.}    \\ha But most people working in quantum gravity will agree that there is some fundamental microstructure to spacetime (``atoms of spacetime'') and the description of spacetime by metric is an approximate long distance description. So why are you making a big deal? I don't see anything novel here.  \\me I will go farther than just saying there is microstructure and show you how to  actually use the thermodynamic concepts to provide an emergent description of gravity --- which no one else has attempted. If you think of the full theory of quantum gravity as analogous to   statistical mechanics then  I will provide   the thermodynamic description of the same system.  As you know, thermodynamics was developed and used effectively decades before we knew anything about the molecular structure of matter or its statistical mechanics. Similarly, even without knowing the microstructure of spacetime or the full quantum theory of gravity, we can make lot of progress with the thermodynamic description of spacetime. The horizon thermodynamics, I will claim, provides \\cite{tworeviews} valuable insights about the nature of gravity totally independent of what ``the atoms of spacetime'' may be. It is somewhat like being able to describe or work with  gases or steam engines without knowing anything about the molecular structure of the gas or steam. ",
        "conclusions": "\\ha We have covered a lot of ground some of which is purely technical while the rest are conceptual or interpretational. In your mind the distinction may be unimportant but others will react differently to results which can be rigorously proved compared to interpretational aspects, however elegant the latter may be. May be you would care to separate them out and provide a summary?  \\me Fine. From a purely algebraic point of view, without bringing in any physical interpretation or motivation, we can prove the following mathematical results: \\begin{itemize} \\item Consider a functional of null vector fields $n^a(x)$ in an arbitrary spacetime given by \\eq{ent-func-2} [or, more generally, by \\eq{generalS}]. Demanding that this functional is an extremum for all null vectors $n^a$ leads to the field equations for the background geometry given by $(2E_{ab}-T_{ab})n^an^b=0$ where $E_{ab}$ is given by \\eq{genEab1} [or, more generally, by \\eq{genEab}]. Thus field equations in a wide class of theories of gravity can be obtained from an extremum principle without varying the metric as a dynamical variable. \\item These field equations are invariant under the transformation $T_{ab} \\to T_{ab} + \\rho_0 g_{ab}$, which relates to the freedom of introducing a  \\cc\\ as an integration constant in the theory. Further, this symmetry forbids the inclusion of a cosmological constant  term in the variational principle by hand as a low energy parameter. \\textit{That is, we have found a symmetry which makes the bulk \\cc\\ decouple from the gravity.} When linearized around flat spacetime, the graviton inherits this symmetry and does not couple to the \\cc . \\item On-shell, the functional in \\eq{ent-func-2} [or, more generally, by \\eq{generalS}] contributes only on the boundary of the region. When the boundary is a horizon, this terms gives precisely the Wald entropy of the theory. \\end{itemize}  \\ha It is remarkable that you can derive not only Einstein's theory uniquely in $D=4$ but even \\LL\\ theory in $D>4$ from an extremum principle involving the null normals \\textit{without varying $g_{ab}$ in an action functional!}. I also see from \\eq{details1} that in the case of Einstein's theory, you have an Lagrangian $n^a(\\nabla_{[a}\\nabla_{b]})n^b$ for a vector field $n^a$ which becomes vacuous in flat spacetime in which covariant derivatives become partial derivatives. There is clearly no dynamics in $n^a$ but they do play a crucial role. So I see that you can get away without ever telling me what the null vectors $n^a$ actually means because they disappear from the scene after serving their purpose. While this may be mathematically clever, it is very unsatisfactory physically. You may not have a rigorous model for these degrees of freedom but what is your picture?  \\me My picture is made of the following ingredients, each  of which seems reasonable but far from having rigorous mathematical justification at this stage.   \\begin{itemize} \\item Assume that the spacetime is endowed with certain microscopic degrees of freedom capable of exhibiting thermal phenomena. This is just the Boltzmann paradigm: \\textit{If one can heat it, it must have microstructure!}; and one can heat up a spacetime. \\item Whenever a class of observers perceive a horizon, they are ``heating up the spacetime'' and the  degrees of freedom close to a horizon  participate in a very \\textit{observer dependent} thermodynamics. Matter which flows close to the horizon (say, within a few Planck lengths of the horizon) transfers energy to these microscopic, near-horizon, degrees of freedom \\textit{as far as the observer who sees the horizon is concerned}. Just as  entropy of a normal system at temperature $T$ will change by $\\delta E/T$ when we transfer to it an energy $\\delta E$, here also an entropy change will occur. (A freely falling observer in the same neighbourhood, of course, will deny all these!) \\item We proved that when the field equations of gravity hold, one can interpret this entropy change in  a purely geometrical manner involving the Noether current. From this point of view, the  normals $n^a$ to local patches of null surfaces are related to the (unknown) degrees of freedom that can participate in the thermal phenomena involving the horizon. \\item Just as demanding the validity of special relativistic laws with respect to all freely falling observers leads to the kinematics of gravity, demanding the local entropy balance in terms of the thermodynamic variables as perceived by local Rindler observers leads to the field equations of gravity in the form $(2E_{ab}-T_{ab})n^an^b=0$.  \\end{itemize}   \\ha It is an interesting picture but is \\textit{totally observer dependent}, right? A local Rindler observer or an observer outside a black hole horizon might attribute all kinds of thermodynamics and entropy changes to the horizons she perceives. But an inertial observer or an observer falling through the Schwarzschild horizon will see none of these phenomena.  \\me Exactly. I claim we need to accept the fact that a whole lot of thermodynamic phenomena needs to be now thought of as observer dependent. For example, if you throw some hot matter on to a Schwarzschild black hole, then when it gets to a few Planck lengths away from the horizon and hovers around it, I expect it to interact with the microscopic horizon degrees of freedom \\textit{as far as an outside observer is concerned}. After all, such an observer would claim that all matter stays arbitrarily close but outside the horizon for all eternity. A freely falling observer through the horizon will have a completely different picture but we have learnt to live with this dichotomy as far as elementary kinematics goes. I think we need to do the same as regards thermodynamics and quantum processes.  \\ha In a way, every key progress in physics involved realizing that something we thought as absolute is  not absolute. With special relativity it was the flow of time and with general relativity it was the concept of global inertial frames and when we brought in quantum fields in curved spacetime it was the notion of particles.  \\me And the notion of temperature, don't forget that. We now know that the temperature attributed to even vacuum state depends on the observer. We need to go further and integrate the entire thermodynamic machinery --- involving highly excited semi-classical states, say, cups of tea with (what we believe to be) ``real'' temperature --- with this notion of LRFs having their own temperature. I don't think this has been done in a satisfactory manner yet \\cite{tpinprogress}.         \\ha So, what next? What is the ``to do list''?  \\me  To name a few which comes to ones mind,  I can list them as  technical ones  and conceptual ones. On the technical side:  (i) It would be nice to make the notion of LRF and the horizon a bit more rigorous. For example, the idea of an approximate Killing vector could be made more precise and one might like to establish the connection between locality, that is apparent in the Euclidean sector and the causality, which is apparent in the Lorentzian section that has the light cones. By and large, one would like to make rigorous the use of LRFs  by, say, computing the next order corrections.  (ii) A lot more can be done to clarify the observer dependence of the entropy. (The study of horizon thermodynamics makes one realize that one does not really quite understand what entropy is!). It essentially involves \\textit{exact} computation of  $(S_1 - S_0)$ where $S_1 $ and $S_0$ are the entropies attributed to the excited and ground state by a Rindler observer. This should throw more light on the expression for $\\delta S$ used in \\eq{delS}. In particular, it would be nice to have a detailed model which shows why $\\delta S$ involves the combination of $\\delta E$ of matter and $T$ of the horizon. One would then use these insights to understand why \\eq{Smatt} actually represents the relevant entropy functional for matter for arbitrary $T_{ab}$.  (iii) It will be nice to have a handle on the positivity or otherwise of the entropy functional used in \\eq{ent-func-2}.  These technical issues, I believe can be tackled in more or less straightforward manner, though the mathematics can be fairly involved. But as I said, they are probably not crucial to the alternative perspective or its further progress. The latter will depend on more serious conceptual issues, some of which are the following:  (i) How come the microstructure of spacetime  exhibits itself indirectly through the horizon temperature even at scales much larger than Planck length? I believe this is because the event horizon works as some kind of magnifying glass allowing us to probe trans-Planckian physics \\cite{magglass} but this notion needs to be made more precise.  (ii) How does one obtain the expression for entropy in \\eq{ent-func-2} from some microscopic model? In particular, such an analysis  --- even with a toy model --- should throw more light on why normals to local patches of null surfaces play such a crucial role as effective degrees of freedom in the long wavelength limit. Of course, such a model should also determine the expression for $P^{abcd}$ and get the metric tensor and spacetime as derived concepts - a fairly tall order!. (This is somewhat like obtaining theory of elasticity starting from a microscopic model for a solid, which, incidentally, is not a simple task either.)     \\ha But what about the ``deep questions'' like, for example, the physics near the singularities? Since you get the same field equations as anybody else does, you will have the same solutions, same singularities etc.  \\me I told you that I am not doing statistical mechanics (which would be the full quantum theory) of spacetime but only thermodynamics. To answer issues related to singularities etc., one actually needs to discover the statistical mechanics underlying the thermodynamic description I have presented here. We can have another chat, after I figure out the statistical mechanics of the spacetime microstructure!"
    },
    "0910/0910.0008_arXiv.txt": {
        "abstract": "We illustrate how  recently improved low-redshift  cosmological measurements can tighten constraints on neutrino properties. In particular we  examine the impact of the assumed cosmological model on the constraints. We first consider the new HST $H_0 = 74.2 \\pm 3.6$ measurement by Riess et al. (2009) and the $\\sigma_8 (\\Omega_m/0.25)^{0.41} = 0.832 \\pm 0.033$ constraint from Rozo et al. (2009) derived from the SDSS maxBCG Cluster Catalog.  In a $\\Lambda$CDM model and when combined with WMAP5 constraints, these low-redshift measurements constrain $\\sum m_{\\nu} < 0.4$ eV at the 95\\% confidence level. This bound does not relax when allowing for the running of the spectral index or for primordial tensor perturbations.  When adding also Supernovae and BAO constraints, we obtain a 95\\% upper limit of $\\sum m_{\\nu}<0.3$eV. We test the sensitivity of the neutrino mass constraint to the assumed expansion history by  both allowing a dark energy equation of state parameter $w\\ne -1$ and by studying a model with coupling between dark energy and dark matter, which allows for variation in $w$, $\\Omega_k$, and dark coupling strength $\\xi$.  When combining CMB,  $H_0$ and the SDSS LRG halo power spectrum from  Reid et al. 2009, we find that in this very general model, $\\sum m_{\\nu} < 0.51$ eV with 95\\% confidence.   If we allow the number of relativistic species $N_{\\rm rel}$ to vary in a $\\Lambda$CDM model with $\\sum m_{\\nu} = 0$, we find $N_{\\rm rel} = 3.76^{+0.63}_{-0.68} (^{+1.38}_{-1.21})$ for the 68\\% and 95\\% confidence intervals.  We also report prior-independent constraints, which are in excellent agreement with the Bayesian constraints. ",
        "introduction": "Atmospheric and solar neutrino experiments have demonstrated that neutrinos have mass, implying a lower limit on the total mass of 0.056 eV \\cite{lesgourgues/pastor:2006}.  Ongoing and future direct experiments will be sensitive to the absolute neutrino mass scale:  for instance, KATRIN will measure or constrain the electron neutrino mass to the 0.2 eV level.\\footnote{{http://www-ik.fzk.de/tritium/motivation/sensitivity.html}}  A thermal neutrino relic component in the universe can impact both the expansion history and growth of structure.  Cosmological probes of neutrinos are therefore complementary to direct experiments and provide some of the tightest constraints on both the number of relativistic species present in the early universe and the sum of neutrino masses \\cite{lesgourgues/pastor:2006}.  However, up to date, these constraints were dependent on the assumed cosmological model.  For example, bounds on the sum of the neutrino masses could be relaxed by, for example, allowing coupling in the dark sector  \\cite{lavacca/bonometto/colombo:2009} or allowing hadronic axions \\cite{hannestad/etal:2007}.  In this paper we  summarize how cosmic neutrinos affect cosmic history, examine how recently improved  cosmological measurements of  the low redshift universe can help constrain these neutrino properties, and scrutinize the impact of our assumptions about the cosmological model.  In the standard model for particle physics  there are three massless neutrino species with weak interactions. Neutrinos decouple early in cosmic history, when the temperature was $T \\sim 1$ MeV, and thereafter contribute to the relativistic energy density with an effective number of species, $N_{\\rm eff} = 3.046$ \\cite{lesgourgues/pastor:2006}. Cosmology is sensitive to the physical energy density in relativistic particles in the early universe $\\omega_{\\rm rel}$ which, in the standard model (for cosmology), includes only photons and neutrinos $\\omega_{\\rm rel}=\\omega_{\\gamma}+N_{\\rm eff}\\omega_{\\nu}$ where $\\omega_{\\gamma}$ is the energy density in photons and $\\omega_{\\nu}$ is the energy density in one active neutrino. As $\\omega_{\\gamma}$ is extremely well constrained, $\\omega_{\\rm rel}$ can be used to constrain neutrino properties;  deviations from  $N_{\\rm eff} = 3.046$  would signal non-standard neutrino features or additional relativistic relics.  $N_{\\rm eff}$ impacts the big bang nucleosynthesis epoch through its effect on the expansion rate.  Therefore measurements of the primordial abundance of light elements can constrain $N_{\\rm eff}$ \\cite{cyburt:2004, cyburt/etal:2005, steigman:2007, iocco/etal:2009}.  These constraints rely on physics at the time of big-bang nucleosynthesis ($T\\sim $ MeV), and in several non-standard models the energy density in relativistic species can change  at later time (e.g. at the last scattering, $T \\sim$ eV). Free-streaming relativistic particles affect the Cosmic Microwave Background (CMB) through their relativistic energy density, which alters the epoch of matter-radiation equality, and through their anisotropic stress \\cite{trotta/melchiorri:2005,komatsu/etal:2009}.  Because the redshift of matter-radiation equality is well constrained by the ratio of the third to first CMB peak height, the first effect defines a degeneracy between $\\Omega_m h^2$ and $N_{\\rm eff}$ \\cite{bowen/etal:2002}: \\begin{equation} \\label{zequal} 1+z_{eq} = \\frac{\\Omega_m h^2}{\\Omega_{\\gamma} h^2}\\frac{1}{1+0.2271 N_{\\rm eff}}. \\end{equation} Because of the second effect, WMAP5 data favors $N_{\\rm eff} = 3.046$ over $N_{\\rm eff} = 0$ at the 99.5\\% confidence level \\cite{dunkley/etal:2009}.  Neutrinos with mass $\\ltap1$ eV become non-relativistic after the epoch of recombination probed by the CMB, so that allowing massive neutrinos alters matter-radiation equality for fixed $\\Omega_m h^2$. Their radiation-like behavior at early times  changes the expansion rate, shifting the peak positions, but this is somewhat degenerate with other cosmological parameters. Therefore, WMAP5 alone constrains $\\sum m_{\\nu} < 1.3$ eV at the 95\\% confidence in a flat $\\Lambda$CDM universe, and  this constraint relaxes to $\\sum m_{\\nu} < 1.5$ eV for a flat $w$CDM universe, assuming $N_{eff} = 3.046$ \\cite{dunkley/etal:2009}.  Figure 17 of \\cite{komatsu/etal:2009} shows that there remain degeneracies between $\\sum m_{\\nu}$ and both $H_0$ and $\\sigma_8$; in this paper we show that currently available low-redshift measurements can break these degeneracies, yielding improved  and robust constraints on $\\sum m_{\\nu}$.  After the neutrinos become non-relativistic, their free-streaming damps power on small scales and therefore modifies the matter power spectrum in the low redshift universe. The latest large-scale structure constraints from the Sloan Digital Sky Survey Luminous Red Galaxy sample in combination with WMAP5 yield $\\sum m_{\\nu} < 0.62$ eV in a flat $\\Lambda$CDM model with $N_{\\rm eff} = 3.046$ \\cite{reid/etal:2009}.  In addition, because the measurement of the matter power spectrum provides a separate constraint on the transfer function, this data combination also yields a constraint on $N_{\\rm eff} = 4.8^{+1.8}_{-1.7}$ in a flat $\\Lambda$CDM cosmology with no massive neutrinos.  We will show that the combination of low redshift data used in this paper provides tighter neutrino constraints in both of these cases.  Relaxing the assumptions about the cosmic expansion history, e.g., requiring $w=-1$ vs. allowing $w$ as a free parameter \\cite{komatsu/etal:2009,Verde/etal:2003} or allowing coupling in the dark sector can relax constraints on neutrino masses \\cite{lavacca/bonometto/colombo:2009}.  As an example we study constraints on neutrino masses in the model of \\cite{gavela/etal:2009}, allowing $\\Omega_k$, $w$, and coupling strength $\\xi$ to vary simultaneously.  In Section~\\ref{whichconstraints} we discuss the low-redshift observational constraints we will make use of to improve neutrino constraints, paying careful attention to how the constraints generalize to the models we consider.  In Sections \\ref{sumneutrinomass} and \\ref{nrelsec} we present new constraints on neutrino masses and the effective number of relativistic species at matter-radiation equality.  We also make a detailed comparison with results from different cosmological probes available in the literature.  In Section~\\ref{conc} we summarize our findings and argue that they are robust to a wide variety of alterations of the cosmological model. ",
        "conclusions": "\\label{conc} In this paper we have demonstrated that probes of the cluster mass function and increased precision on the Hubble constant can break key degeneracies with CMB observations and yield excellent constraints on both the number of species and sum of the masses of cosmological neutrinos.  When the expansion history is fixed to $\\Lambda$CDM, current constraints on $H_0$ \\cite{riess/etal:2009} and the cluster mass function \\cite{rozo/etal:2009} constrain $\\sum m_{\\nu} < 0.4$ eV at 95\\% confidence.  This bound relaxes to 0.5 eV in the two extended models we considered: wCDM, and the dark coupling model of \\cite{gavela/etal:2009}, which allowed curvature, $w$, and coupling strength $\\xi$ to vary.  Probing the mass function using X-ray clusters, Ref.~\\cite{vikhlinin/etal:2009} combines X-ray cluster data with WMAP5+BAO+SN and finds $\\sum m_{\\nu} < 0.33$ eV in a $w$CDM cosmology, though the systematic errors still must be clearly quantified.  The optimistic upper bound available from current data, $\\sum m_{\\nu} \\sim 0.3$ eV, almost excludes the scenario in which the neutrino masses are quasi-degenerate.  The constraint on the number of relativistic species $N_{\\rm rel}$ from probes of the cluster mass function has not been  widely explored so far.  We find that the combination of WMAP5, \\cite{riess/etal:2009} $H_0$, and the maxBCG cluster mass function constraint provides an excellent constraint: $N_{\\rm rel} = 3.76^{+0.63}_{-0.68}$.  However, we point out that this constraint does not improve when BAO and SN data are also included; those data sets have been shown to have excellent constraining power on both $\\Omega_k$ and $w$ \\cite{percival/etal:2009}.  Therefore, we may expect that the constraint on $N_{\\rm rel}$ will not relax dramatically if we allow for these additional parameters but include all of these datasets, particularly because of the complementarity of the growth of structure and expansion history constraints we have employed.  These constraints are not competitive with the most optimistic errors from BBN \\cite{iocco/etal:2009}, but provide an important consistency check using probes significantly after the BBN epoch."
    },
    "0910/0910.2461_arXiv.txt": {
        "abstract": "{PAHs appear to be an ubiquitous interstellar dust component but the effects of shocks waves upon them have never been fully investigated.} {We study the effects of energetic ($\\approx 0.01-1$ keV) ion (H, He and C) and electron collisions on PAHs in interstellar shock waves.} {We calculate the ion-PAH and electron-PAH nuclear and electronic interactions, above the threshold for carbon atom loss from a PAH, in $50-200$ km s$^{-1}$ shock waves in the warm intercloud medium.} {Interstellar PAHs ($N_{\\rm C} = 50$) do not survive in shocks with velocities greater than 100 km s$^{-1}$ and larger PAHs ($N_{\\rm C} = 200$) are destroyed for shocks with velocities $\\geq 125$ km s$^{-1}$. For shocks in the $\\approx 75 - 100$ km s$^{-1}$ range, where destruction is not complete, the PAH structure is likely to be severely denatured by the loss of an important fraction ($20-40$\\%) of the carbon atoms. We derive typical PAH lifetimes of the order of a few $\\times 10^8$ yr for the Galaxy. These results are robust and independent of the uncertainties in some key parameters that have yet to be well-determined experimentally.} {The observation of PAH emission in shock regions implies that that emission either arises outside the shocked region or that those regions entrain denser clumps that, unless they are completely ablated and eroded in the shocked gas, allow dust and PAHs to survive in extreme environments.} ",
        "introduction": "Interstellar Polycyclic Aromatic Hydrocarbon molecules (PAHs) are an ubiquitous component of the interstellar medium. The mid-infrared spectrum of the general diffuse interstellar medium as well as energetic environments near massive stars such as {\\sc H\\,ii} regions and reflection nebulae are dominated by broad emission features at 3.3, 6.2, 7.7, and 11.2 $\\mu$m. These emission features are now generally attributed to infrared fluorescence by large PAH molecules containing 50-100 C-atoms, pumped by single FUV photons \\citep[see] [ for a recent review]{tielens08}. The observed spectra also show evidence for PAH clusters containing a few hundred C-atoms \\citep{breg89, rap05, berne07} as well as very small dust grains \\citep[$\\sim$30 \\AA; ][]{desert90}. It seems that the interstellar grain size distribution extends all the way into the molecular domain \\citep{allamandola89, desert90, draine01}. The origin and evolution of interstellar PAHs are somewhat controversial. On the one hand, based upon extensive laboratory studies of soot formation in terrestrial environments, detailed models have been made for the formation of PAHs in the ejecta of C-rich giants \\citep{frenk89, cher92} -- as intermediaries or as side-products of the soot-formation process -- and studies have suggested that such objects might produce enough PAHs to seed the ISM \\citep{latter91}. On the other hand, models have been developed where PAHs (as well as very small grains) are the byproduct of the grinding-down process of large carbonaceous grains in strong supernova shock waves which permeate the interstellar medium \\citep{bork95, jones96}. Grain-grain collisions shatter fast moving dust grains into small fragments and, for graphitic progenitor grains, these fragments might be more properly considered PAH molecules. The destruction of interstellar PAHs is equally clouded. Laboratory studies have shown that small (less than 16 C-atoms), (catacondensed) PAHs are rapidly photodissociated by $\\sim$ 10 eV photons \\citep{joch94}. However, this process is strongly size-dependent as larger PAHs have many more modes over which the internal energy can be divided and PAHs as large as 50 C-atoms might actually be stable against photodissociation in the ISM \\citep{lePage01, allamandola89}. While strong shock waves have been considered as formation sites for interstellar PAHs, the destruction of these PAHs in the hot postshock gas has not been evaluated. Yet, high energy ($\\sim 1$ keV) collisions of PAHs with ions and electrons are highly destructive.  The observational evidence for PAHs in shocked regions is quite ambiguous. The majority of supernova remnants does not show PAH features \\citep[e.g. \\object{Cas A},][]{smith08}, but observations of \\object{N132D} \\citep{tappe06} suggest the possibility of PAH survival in shocks. Recent work by \\citet{and07} investigates the presence of PAHs in a subset of galactic supernova remnants detected in the GLIMPSE survey. Unfortunately the interpretation of such observations is not straightforward, because of the difficulty in disentangling the PAH features intrinsic to the shocked region with those arising from the surrounding material. Another interesting case is the starburst galaxy \\object{M82}, which shows above and below the galactic plane a huge bipolar outflow of shock-heated gas interwoven with PAH emission\\footnote{http://chandra.harvard.edu/photo/2006/m82/} \\citep{armus07}. PAHs have also been observed at high galactic latitudes in the edge-on galaxies \\object{NGC 5907} and \\object{NGC 5529} \\citep{irwin06, irwin07}. Shock driven winds and supernovae can create a so-called \"galactic fountain\" \\citep{bregman80} transporting material into the halo and these detections of PAHs suggests the possibility of survival or formation of the molecules under those conditions. On the other hand \\citet{ohalloran06, ohalloran08} have found a strong anti-correlation between the ratio [FeII]/[NeII] and PAH strenght in a sample of low-metallicity starburst galaxies. Since [FeII] has been linked primarly to supernova shocks, the authors attributed the observed trend to an enhanced supernova activity which led to PAH destruction.  In our previous study \\citep{jones96}, we considered the dynamics and processing of small carbon grains with $N_{\\rm C} \\geq$~100. The processing of these grains by sputtering (inertial and thermal) in ion-grain collisions and by vaporisation and shattering in grain-grain collisions was taken into account for all the considered grain sizes. In that work, the smallest fragments ($a_{\\rm f} < 5$\\,\\AA) were collected in the smallest size bin and not processed. In this work we now consider what happens to these smallest carbon grain fragments that we will here consider as PAHs. In this paper, we will consider relatively low velocity ($\\leq 200$ km s$^{-1}$) shocks where the gas cools rapidly behind the shock front but, because of their inertia, PAHs (and grains) will have high velocity collisions even at large postshock column densities. Collisions between PAHs and the gas ions occur then at the PAH velocity which will slowly decrease behind the shock front due to the gas drag. This relative velocity is thus independent of the ion mass and, for dust grains, destruction is commonly called inertial sputtering. Destruction of PAHs in high velocity ($\\geq 200$ km s$^{-1}$) shocks -- which cool slowly through adiabatic expansion -- is dominated by thermal sputtering and these shocks are considered in a subsequent paper \\citep[][ hereafter MJT]{micelotta09}.  This paper is organized as follows: Sect. 2 describes the theory of ion interaction with solids, Sect. 3 illustrates the application of this theory to PAH processing by shocks and Sect. 4 presents our results on PAH destruction. The PAH lifetime in shocks and the astrophysical implications are discussed in Sect. 5 and our conclusions summarized in Sect. 6. ",
        "conclusions": "We have extensively studied the effects of PAH processing by shocks with velocities between 50 and 200 km s$^{-1}$, in terms of collisions with ions and electrons which can lead to carbon atom loss, with a consequent disruption and destruction of the molecule.  An ionic collision consists of two simultaneous processes which can be treated separately: a binary collisions between the projectile and one single atom in the target (nuclear interaction) and the energy loss to the atomic electrons (electronic interaction). For the nuclear interaction, we modified the existing theory in order to treat collisions able to transfer energy \\textit{above} a specific threshold. This is the case we are interested in, which has not been treated in previous studies. For electronic interaction and collisions with electrons we developed specific models for PAH, described in MJT.  The PAH dynamics in the shocks is evaluated using the same approach as in our previous work \\citep{jones94, jones96}. For nuclear interaction, the level of carbon atom loss increases for decreasing values of the threshold energy $T_{0}$. We adopt $T_{0}$ = 7.5 eV as a reasonable value, but experimental determinations are necessary. In ionic collisions, the carbon contribution to PAH destruction is totally negligible because of its very low abundance with respect to H and He. The fractional destruction induced by nuclear excitation increases with the PAH size, while in case of electronic excitation and electron collisions is lower for higher $N_{\\rm C}$ values, i.e. bigger PAHs are more resistant than the smaller ones against electron and electronic processing.  The parameter $E_{0}$ required for the evaluation of PAH destruction due to electron and electronic interaction is unfortunately not well constrained. We adopt a value of 4.58 eV consistent with extrapolations to interstellar conditions, but better determinations would be desirable.  Electronic interaction, both inertial and thermal, plays a marginal role in PAH processing by shocks. We find that 50 carbon atoms PAHs are significantly disrupted in ionic collisions for shock velocities below 75 km s$^{-1}$, mainly by inertial `sputtering' by helium during nuclear interaction. Our results indicate $5-15\\,\\% $ C atom loss, sufficient to cause a severe de-naturation of the PAH aromatic structure. Above 100~km~s$^{-1}$ such PAHs are instead totally destroyed by collisions with thermal electrons. For $N_{\\rm C} = 200$, PAHs experience increasing damaging caused by nuclear ionic collisions up to 100~km~s$^{-1}$, which turns into complete atomic loss for higher velocities.  In this case the destruction is due to the combined effect of electrons and nuclear interaction with thermal ions.  The calculated PAH lifetime against destruction, $t_{\\rm SNR}$, is 1.6 $\\times$10$^{8}$ yr and 1.4 $\\times$10$^{8}$ yr for $N_{\\rm C}$ = 50 and 200 respectively. Small PAHs are preferentially destroyed by electrons, big PAHs by ions. The calculated lifetimes are smaller than the values found for carbonaceous grains (6 $\\times$10$^{8}$ yr) but close to that for hydrogenated amorphous carbon ($~$2 $\\times$10$^{8}$ yr), and far from the stardust injection timescale of 2.5 $\\times$10$^{9}$ yr. The presence of PAHs in shocked regions therefore requires an efficient reformation mechanism and/or a protective environment.  We surmise that the molecular structure of PAHs is strongly affected by shock processing in the interstellar medium. Electronic excitation by impacting ions or electrons may lead to isomerization into stable pure-C species such as fullerenes. In contrast, nuclear interaction may lead to the formation of N-containing PAHs. Further laboratory studies are required to demonstrate the viability of these chemical routes.   \\appendix"
    },
    "0910/0910.4184_arXiv.txt": {
        "abstract": "We report studies of ultra-high energy cosmic ray composition via analysis of depth of airshower maximum ($\\xmax$), for airshower events collected by the High Resolution Fly's Eye (HiRes) observatory. The HiRes data are consistent with a constant elongation rate $d\\xmean/d(log(E))$ of $47.9 \\pm 6.0 ({\\rm stat.}) \\pm 3.2 ({\\rm syst.})$~g/cm$^2$/decade for energies between 1.6~EeV and 63~EeV, and are consistent with a predominantly protonic composition of cosmic rays when interpreted via the QGSJET01 and QGSJET-II high-energy hadronic interaction models. These measurements constrain models in which the galactic-to-extragalactic transition is the cause of the energy spectrum ``ankle'' at $4 \\times 10^{18}$~eV. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.1597_arXiv.txt": {
        "abstract": "Normal modes of oscillation of the Sun are useful probes of the solar interior.  In this work, we use the even-order splitting coefficients to study the evolution of magnetic fields in the convection zone over solar cycle 23, assuming that the frequency splitting is only due to rotation and a large scale magnetic field.  We find that the data are best fit by a combination of a poloidal field and a double-peaked near-surface toroidal field. The toroidal fields are centered at $r_0=0.999R_\\odot$ and $r=0.996R_\\odot$ and are confined to the near-surface layers.  The poloidal field is a dipole field.  The peak strength of the poloidal field is $124\\pm17$~G. The toroidal field peaks at $380\\pm30$~G and $1.4\\pm 0.2$~kG for the shallower and deeper fields respectively.  The field strengths are highly correlated with surface activity.  The toroidal field strength shows a hysteresis-like effect when compared to the global 10.7 cm radio flux.  The poloidal field strength shows evidence of saturation at high activity. ",
        "introduction": "Understanding the nature of the Sun's magnetic fields --- their structure and variability, their generation mechanisms, and their effects on the heliosphere --- is one of the key aims of current research in solar physics.  It is generally believed that the magnetic fields are generated by a cyclic dynamo that operates somewhere in the solar interior.  In this paper, we use helioseismology to study the global scale internal magnetic fields over the course of solar cycle 23.  Helioseismology is the most powerful tool available to solar physicists to study the interior of the Sun.  The oscillation frequencies have been used to study the structure and dynamics of the solar interior with great precision. Magnetic fields, however, have proved to be much more challenging.  There are a number of important difficulties in dealing with magnetic fields in a helioseismic context.  The magnitudes of the signatures in the data are quite small, making statistically significant measurements challenging.  Secondly, the interpretation of data is very difficult.  The physics of wave propagation in the presence of magnetic fields is far more complex than in the non-magnetic case.  Further, the geometry of the underlying field strongly affects the signatures in helioseismic global mode frequencies, meaning different field configurations and strengths can be difficult to distinguish from their helioseismic signatures.  Even worse, \\citet{ZG95} showed that because magnetic fields act on mode frequencies both by perturbing the thermal structure of the Sun and by changing the wave propagation speeds directly, there is a degeneracy between magnetic field effects and other thermal perturbations which cannot be distinguished a priori from helioseismic data.  Although helioseismic determinations of magnetic fields are difficult, there have been many attempts to do so.  \\citet{Isaak82} suggested that the then observed frequency splittings in the solar acoustic spectrum could be caused by a large scale magnetic field situated in the core. \\citet{DG84} used an asymptotic approximation to study the effects of magnetic field on the splitting coefficients, and \\citet{DG88} argued that a 1 MegaGauss field at the base of the convection zone was necessary to explain the observed splitting coefficients.  However, \\citet{Basu97, ACT00} placed a limit of 0.3~MG on the field at the base of the convection zone; thus the situation was unclear.  A mega-Gauss magnetic field is also inconsistent with dynamo theories and constraints from other observations \\citep[e.g.,][]{DC93}.  \\citet{GT90} developed a formalism to compute the effects of rotation and axisymmetric magnetic fields on the frequency splittings (discussed in the following section), which \\citet{ACT00} used to analyze the first year of Michelson Doppler Imager (MDI) data.  They placed limits on the strengths of internal toroidal fields, finding a limit of 20~kG at a depth of 30~Mm, and a limit of 300~kG at the base of the convection zone ($r = 0.713R_\\odot$). \\citet{DGKS00} inverted the mean frequencies and splitting coefficients for changes in temperature, and found that the resulting temperature perturbation could be explained by a change in magnetic field of 60~kG at a depth of 45~Mm ($r\\sim0.93R_\\odot$).  \\citet{DGS01} found that changes in $f$-mode frequencies from solar minimum to solar maximum implied a decrease in solar radius with activity, which they associated with a change between 4 and 8~Mm in depth.  In explaining this result with changing magnetic fields, they assumed a tangled field, but even so the magnitude of the change in field strength was strongly dependent on the radial distribution of the field.  The change they required was 7~kG for a uniform field, or substantially less (1~kG at 8~Mm) for an inwardly increasing field.  \\citet{CS02,CS05} looked for signatures of a change at the base of the convection zone from low activity to high activity, and found signs of a small change, which they proposed could be due to a change in magnetic field of 170 -- 290~kG. \\citet{BB08}, working with an entire solar cycle's worth of helioseismic data, found a change in sound speed between solar maximum and solar minimum at the base of the convection zone, which, if due to a change in magnetic field, could indicate a change in field strength of 290~kG at that depth.  In this work, we exploit the fact that we have much  more helioseismic data than previous investigators had access to, and try to get a coherent picture of sub-surface solar magnetic fields and their temporal evolution.  We extend the work of \\citet{ACT00}, who considered toroidal magnetic fields, to include poloidal fields.  This means that we can, in principle, consider any axisymmetric magnetic field configuration.  We compute the effects of a wide variety of magnetic field configurations on the $a_2$ splitting coefficients, and compare them to a solar cycle's worth of MDI data. It is not clear if the solar magnetic field has large scale structure of the form we assume or whether it is in tangled state due to turbulence in the convection zone. Since the effect of magnetic field manifests through a quadratic term in magnetic field, our estimate may also be applicable to tangled field with some degree of approximation. ",
        "conclusions": "We have attempted to use the first even order splitting coefficient ($a_2$) to infer the configuration and strength of the Sun's internal magnetic fields over the course of solar cycle 23, assuming that the entire signature in $a_2$ after correction for rotation effects is magnetic and that the fields are axisymmetric.  The field that we have found is a combination of poloidal field and a double-peaked near-surface toroidal field. The strengths of the poloidal and toroidal components, at least for high activity period, are well correlated.  The relative strengths of the two toroidal fields are also extremely well correlated.  Although the fits we have shown are the best fit to the data from the grid of models that we have computed, we can say nothing about the uniqueness of these fits over the set of all possible magnetic field configurations in the solar interior.  In particular, the choice of radial profile of the toroidal fields is virtually limitless, and by restricting our work to profiles of the form (\\ref{tor}), we have limited our search to a restricted class of fields. It is possible that there are fields we did not consider with quite different radial and latitudinal distributions which fit the data as well as the fields we have presented as best fits.  In addition, as noted above, it is not strictly correct to add the splitting coefficient perturbations together as we have done without explicitly accounting for the perturbations arising from the cross terms.  We do not expect, however, these corrections to be significant, and a full treatment of these corrections would be considered in a future work.  Our inferred magnetic field does not change its latitudinal distribution over the course of the solar cycle.  This is in part due to the fact that we are only fitting the $a_2$ coefficient (as noted above, the higher order splittings had large errors), so our sampling of the interior is not really latitudinally sensitive.  Thus, we do not see a butterfly diagram in our magnetic fields.  \\citet{UB05} measured the surface toroidal component of the solar magnetic field over almost an entire 22-year cycle.  The field they measure is roughly a tenth of the peak strength of our toroidal field.  Strictly speaking, however, we see no toroidal field at all at the surface, since in our inferred field, the field strength becomes zero precisely at $r=1R_\\odot$. The peak strength that we measure, however, is only 700~km below the surface, and the field could penetrate the surface somewhat.  \\citet{UB05} find a field which gives a $\\beta \\sim 6 \\times 10^{-5}$ at the surface at high activity, and drops to nothing at low activity, while we find a field that changes from $\\beta_0 \\sim 10^{-4}$ at low activity to $\\beta_0=2\\times10^{-3}$ at high activity at a radius of $r=0.999R_\\odot$ (a depth of approximately 700~km).  Recently, attention has been focused on the strength and configuration of the quiet Sun surface magnetic field.  \\citet{Harveyetal07} reported the presence of a `seething' horizontal magnetic field with an rms field strength of 1.7~G.  With the launch of Hinode (Solar B), the high spatial resolution of the onboard spectropolarimeter has been used to study the horizontal fields of the solar photosphere.  \\citet{Lites07,Lites08} have measured the horizontal flux, which they find to be 55~G, compared to the average vertical flux of 11~G.  \\citet{PP09} found that the zonal component (component in the East-West direction) was much smaller than the radial component, reporting an inclination angle of less than $12^\\circ$ from vertical in the East-West direction.  The fields being studied in the aforementioned works are generally very tangled fields which thread through the intergranular lanes and so they are not axisymmetric fields.  It is worthwhile to compare our results with theirs, since tangled fields on local scales can organize into roughly axisymmetric fields on global scales. However, the contribution to splittings are more sensitive to $\\langle B^2\\rangle$ rather than $\\langle B\\rangle^2$ and hence tangled field may also contribute to it, even when the average $\\langle B\\rangle$ is very small. Further, considering the general behavior of the perturbation to the mode frequencies, our inference about the location of required magnetic field is more robust as a different location will yield a very different behavior of splitting coefficients. The exact magnitude of the field may depend on the assumption of geometry and on it being tangled or large scale. Nevertheless, we believe that our estimate is of the right order, though the statistical errorbars obtained by us may not be realistic. The systematic errors in these estimates would be certainly larger.  The dominance of poloidal field orientation at the surface found by \\citet{PP09} is found in our own results --- at the surface, the toroidal field is weak or vanishing, but the poloidal field remains.  In the period analyzed by \\citet{Lites07,Lites08}, we find a poloidal field strength of 40~G, and a toroidal field of 90~G at a depth of 700~km.  The vertical flux they find (11~G) is weaker than what we detect, but their 55~G horizontal flux may be roughly consistent with our toroidal field.  \\citet{SL08} found that the dipole moment of the surface magnetic field, measured from MDI magnetograms, was half the strength in 2008 that it was in 1997, during the last solar minimum.  We do not see such a difference from the beginning of our period to the end --- in fact, we find the poloidal field strength is slightly higher during the current minimum, although the level of the difference is within the errors, and our data sets end in 2007, so the comparison is not contemporaneous.  Hysteresis in the relations between activity indices has been observed before, for example in the relation between low degree \\citep{AGetal92,J-Retal98} and intermediate degree \\citep{Tripathyetal00,Tripathyetal01} acoustic modes and global magnetic indices.  It should be noted that an analysis of a full solar cycle's worth of intermediate degree $p$-modes data does not show any hysteresis in mean frequencies as a function of 10.7~cm flux \\citep{BB08}. \\citet{Tripathyetal01} noted that, among the global mode indices, the relation between global line-of-sight magnetic flux and 10.7~cm radio flux showed a hysteresis effect, but the relation between the radiative indices and 10.7~cm flux did not.  \\citet{M-IS00} argued that the observed hysteresis could be almost entirely due to the latitudinal distribution of magnetic flux on the surface of the Sun.  We believe that this is a compelling explanation for the hysteresis that we find.  The 10.7~cm flux is the integrated flux received at the Earth and does not contain any information about the latitudinal variation, while the $a_2$ splitting coefficient is associated with definite latitudinal variation, given by $P_2(\\cos\\theta)$ \\citep{ABH01}, and hence the two would not be the same.  More importantly, we expect the actual magnetic fields in the near surface layers to drift equatorward --- as the surface fields do.  Few conclusions can really be drawn from this work with respect to dynamo theory since the fields we have inferred are predominantly shallow fields, whereas most dynamo models operate much deeper down, in the shear layer at and below the base of the convection zone.  \\citep[some useful recent reviews include][]{Oss2003,Char2005,MT2009}.  The upper limits that we place on fields at that depth are consistent with earlier helioseismic results \\citep[e.g.,][]{Basu97, ACT00, CS02,CS05,BB08}.  Many deep-seated dynamo mechanisms predict an anticorrelation between the poloidal and toroidal field components, as the dynamo converts poloidal to toroidal field and toroidal field back to poloidal.  We do not see any evidence of such conversion.  Some dynamo mechanisms, however, operate in the near-surface shear layer \\citep[e.g.][]{Brand05}.  Although the fields generated in these models are generally extremely tangled, on global scales these fields can have toroidal and poloidal components \\citep[e.g.][]{BBBMNT07,BBMBT09}. In particular, although they were considering a more rapidly rotating star than the Sun, \\citet{BBBMNT07} noted that their field contained both a poloidal and a toroidal component, and that the toroidal component was much the stronger of the two."
    },
    "0910/0910.2288_arXiv.txt": {
        "abstract": "We present observations of interstellar rubidium toward $o$ Per, $\\zeta$~Per, AE Aur, HD 147889, $\\chi$ Oph, $\\zeta$ Oph, and 20 Aql. Theory suggests that stable $^{85}$Rb and long-lived $^{87}$Rb are produced predominantly by high-mass stars, through a combination of the weak s- and r-processes. The $^{85}$Rb/$^{87}$Rb ratio was determined from measurements of the Rb~{\\small I} line at 7800 \\AA\\ and was compared to the solar system meteoritic ratio of 2.59. Within 1-$\\sigma$ uncertainties all directions except HD 147889 have Rb isotope ratios consistent with the solar system value. The ratio toward HD 147889 is much lower than the meteoritic value and similar to that toward $\\rho$ Oph A \\citep{fed04}; both lines of sight probe the Rho Ophiuchus Molecular Cloud. The earlier result was attributed to a deficit of r-processed $^{85}$Rb. Our larger sample suggests instead that $^{87}$Rb is enhanced in these two lines of sight. When the total elemental abundance of Rb is compared to the K elemental abundance, the interstellar Rb/K ratio is significantly lower than the meteoritic ratio for all the sight lines in this study. Available interstellar samples for other s- and r- process elements are used to help interpret these results. ",
        "introduction": "Rubidium isotopes $^{85}$Rb and $^{87}$Rb are products of the neutron capture slow (s-) and rapid (r-) processes. ($^{87}$Rb is unstable with the long halflife of 4.75 $\\times$ 10$^{10}$ years.) The s-process contributions come from two types of stars: the so-called main s-process operates in low mass asymptotic giant branch (AGB) stars with their products ejected into the interstellar medium (ISM) by the stellar wind; and the weak s-process occurs in the He- and C-burning shells of massive stars and distributed into the ISM primarily, it is thought, by a star's terminal supernova explosion. The r-process most likely occurs at the time of the supernova explosion. Of primary relevance to studies of the Rb elemental and isotopic abundances beyond the solar system is the division of responsbility for Rb synthesis between low-mass long-lived and high-mass short-lived stars or, equivalently, the relative contributions from the main s-process and the combination of the weak s-process and the r-process.  Dissection of the solar system abundances for Rb and adjacent elements enables fractional contributions of the weak and main s-processes to be estimated. Then, the r-process contribution is obtained as the differences between the solar system abundance and the sum of the weak and main s-process contributions. The most recent exercise of this kind is that reported by \\citet{heil} who estimated the fractional contributions of the main s-process to the solar system abundances of $^{85}$Rb and $^{87}$Rb to be 0.17 and 0.24, respectively. In other words, the relative contributions of low mass and high mass stars is 0.17/0.83 = 0.20 for $^{85}$Rb and 0.24/0.76 = 0.32 for $^{87}$Rb. [\\citet{heil} give the weak s-process fractional contributions as 0.24 for $^{85}$Rb and 0.46 for $^{87}$Rb and, then by subtraction of the two s-process contributions, the r-process contributions are 0.59 for $^{85}$Rb and 0.30 for $^{87}$Rb.] Earlier dissections of the solar system abundances have given somewhat different results for the relative contributions of low and high mass stars, e.g., \\citet{arland} gave 0.16/0.84 = 0.19 for $^{85}$Rb and 0.35/0.65 = 0.54 for $^{87}$Rb indicating a more pronounced role of the massive stars in the synthesis of $^{87}$Rb (relative to $^{85}$Rb) than suggested by \\citet{heil}. Most probably, higher weight should be given to Heil et al.'s result because it was based in part on new and accurate measurements of neutron capture cross sections, especially for the two Rb isotopes. Our study of interstellar Rb was undertaken to set observational constraints on the proposed mechanisms of Rb nucleosynthesis.  Rubidium has been analyzed in both unevolved and evolved stars. \\citet{lluck} sought the Rb isotope ratio in Arcturus. \\citet{gratton} studied the abundances of neutron-rich elements in metal-poor stars, while \\citet{tomkin} considered rubidium in metal-deficient disk and halo stars. Others have examined rubidium abundances in AGB stars \\citep{lambert,abia2,abia1}. However, rubidium is difficult to measure in stellar atmospheres. The Rb~{\\small I} lines at 7800 and 7947 \\AA\\ are blended with nearby lines and stellar broadening prevents the hyperfine and isotopic components from being resolved. Intrinsically narrow interstellar lines are not hindered by these complications.  The detection of rubidium in the ISM was first reported by \\citet{jura} from absorption toward the bright star $\\zeta$ Oph. They used the Rb~{\\small I} 5{\\it s} $^{2}${\\it S}$_{1/2}$--5{\\it p}~$^{2}${\\it P}$_{3/2}$ line at 7800.29 \\AA. This is the strongest resonance line of neutral rubidium and it is used for our study as well, even though most rubidium in diffuse interstellar clouds is ionized because it has a low ionization potential of 4.177 eV. Rubidium's overall low abundance requires the use of its strongest line.  The early detection by \\citet{jura} could not be confirmed by \\citet{fed85}, who also could not detect Rb~{\\small I} toward $o$ Per nor $\\zeta$ Per. \\citet{fed85} calculated a typical 3-$\\sigma$ upper limit in column density of $\\sim$ 5 $\\times$ 10$^{9}$ cm$^{-2}$ for these lines of sight. Interstellar absorption toward these three stars is revisited in this study. More recently, \\citet{gredel} provided the first firm detections of interstellar rubidium toward Cyg OB2 No. 12 and No. 5.  The solar system isotope ratio is found from carbonaceous chondrite meteorites, with a ratio of $^{85}$Rb/$^{87}$Rb = 2.59 \\citep{lodders}. \\citet{fed04} found an interstellar isotope ratio toward $\\rho$ Oph A of 1.21, which differs significantly from the meteoritic value. This was the first relatively precise determination of the rubidium isotope ratio for extrasolar gas.  In a recent paper, Kawanomoto et al. (2009) studied the line of sight toward HD~169454 and derived an isotope ratio that is consistent with the solar system value. The present study expands upon the results of \\citet{fed04} and Kawanomoto et al. (2009) on the stable $^{85}$Rb and long-lived $^{87}$Rb isotopes.  We sought Rb absorption along seven lines of sight toward $o$ Per, $\\zeta$ Per, AE Aur, HD 147889, $\\chi$ Oph, $\\zeta$ Oph, and 20 Aql. These lines of sight were selected because they sample different directions in the solar neighborhood, and because other relevant data noted below are available for them. Furthermore, the cloud structure is simple with only one or two apparent components per line of sight. We explore $^{85}$Rb/$^{87}$Rb ratios and compare the results to those for $\\rho$~Oph~A. We also compare the elemental abundance ratio of Rb/K, as well as the $^{85}$Rb/K and $^{87}$Rb/K ratios, to examine whether the $^{85}$Rb or $^{87}$Rb abundances are lower or higher relative to the abundances seen in meteorites. Section 2 summarizes the high-resolution observations and data reduction techniques on the Rb~{\\small I} $\\lambda$7800 line. Section 3 presents the analysis of the data and Section 4 the results, including confidence limits on the column densities. Section 5 uses the results to discuss the implications for interstellar rubidium and its nucleosynthetic production site(s). Finally, Section 6 gives a summary and conclusion. ",
        "conclusions": "The interstellar $^{85}$Rb/$^{87}$Rb ratio was measured along seven lines of sight, $o$ Per, $\\zeta$ Per, AE Aur, HD 147889, $\\chi$ Oph, $\\zeta$ Oph, and 20 Aql. Using K~{\\small I} absorption as a guide, the spectra were synthesized to find the velocity, column density, and {\\it b}-value of each cloud. Most clouds had ratios that agreed, within mutual uncertainties, with the solar system value of 2.59. The gas toward HD 147889 had a lower $^{85}$Rb/$^{87}$Rb ratio (1.03~$\\pm$~0.21) and a higher elemental ratio (8.6~$\\pm$~1.3), suggesting that $^{87}$Rb is enhanced in this direction, consistent with the previous results for $\\rho$ Oph A \\citep{fed04}. The anomalous ratios appear to arise within the Rho Ophiuchus Molecular Cloud. Furthermore, a comparison of interstellar Rb to K shows that the gas in the solar neighborhood appears to be significantly underabundant in rubidium, even after considering uncertainties in ionization balance and level of depletion. An examination of other species, such as Ge, Kr, Cd, and Sn, seems to suggest nuclides synthesized by the n-capture processes occuring in massive stars (weak s and r) have lower interstellar abundances than expected. Larger samples for several n-capture elements, from both observation and theory, will help clarify the cause for this."
    },
    "0910/0910.3184.txt": {
        "abstract": "Pseudo-Newtonian gravitational potential introduced in spherically symmetric black-hole spacetimes with a repulsive cosmological constant is tested for equilibrium toroidal configurations of barotropic perfect fluid orbiting the black holes. Shapes and potential depths are determined for the marginally stable barotropic tori with uniform distribution of specific angular momentum, using both the pseudo-Newtonian and fully relativistic approach. For the adiabatic (isoentropic) perfect fluid, temperature profiles, mass-density and pressure profiles and total masses of pseudo-Newtonian and relativistic tori are compared providing important information on the relevance of the test-disc approximation in both the approaches. It is shown that the pseudo-Newtonian approach can be precise enough and useful for the modelling of accretion discs in the Schwarzschild-de Sitter spacetimes with the cosmological parameter $y=\\Lambda M^2/3\\lesssim 10^{-6}$. For astrophysically relevant black holes with $y<10^{-25}$, this statement is tested and shown to be precise in few percent for both accretion and excretion tori and for the marginally bound, i.e., maximally extended tori allowing simultaneous inflow to the black hole and outflow to the outer space. ",
        "introduction": "\\label{intro} Cosmological observations of distant Ia-type supernova explosions indicate an~accelerating universe. Starting at the cosmological redshift $z\\approx 1$, the accelerated expansion should be generated by some appropriate form of the so-called dark energy \\cite{Per-etal:1999:ASTRJ2:,Rie-etal:2004:ASTRJ2:}. These results are in accord with a large variety of cosmological tests including gravitational lensing, galaxy number counts, etc. \\cite{Ost-Ste:1995:NATURE:}. The recent detailed studies of the cosmic microwave background (CMB) anisotropies indicate that the energy content of the dark energy represents $\\sim 74.5\\%$ of the energy content in the observable universe, and the sum of energy densities is very close to the critical energy density $\\rho_{\\rm crit}$, corresponding to almost flat universe \\cite{Spe-etal:2003:ASTJS:,Spe-etal:2007:ASTJS:}.  A large variety of possible candidates for the dark energy is discussed in these days. First of all, there is the standard possibility represented by the cosmological constant $\\Lambda$. Its Lorentz invariant form enables interpretation in terms of a~ground state or vacuum energy of quantum fields \\cite{Kol-Tur:1990:FrontiersinPhysics:,Dol-Zel-Saz:1988:IzdatMoskUniv:}. The energy density $\\rho_{\\Lambda}$, which can be associated with the cosmological constant, remains unchanged during the cosmic expansion, and its pressure to energy density ratio (equation of state) is $w=p_{\\Lambda}/\\rho_{\\Lambda}=-1$. Further, there is a variety of scalar fields evolving outside of their energy minimum, called quintessence, which possess a time varying energy density and equation of state with $w<-1/3$ \\cite{Zla-Wan-Ste:1999:PHYRL:}. Such a scenario can be realised by light scalar field coming from modified $f(R)$ gravity \\cite{Noj-Odi:2003:PHYSR4:}, string-inspired cosmologies \\cite{Tsu-Sam:2007:JOCAP:}, cosmology with extra dimensions \\cite{Neu:2004:CLAQG:}, or by $k$-essence being a scalar field with a non-canonical kinetic term \\cite{Arm-Dam-Muk:1999:PHYLB:,Gar-Muk:1999:PHYLB:}. Similar behaviour is exhibited by the coupled dark energy, i.e., a scalar field coupled to the dark matter. For example, Chaplygin gas or its generalization called quartessence explain both dark energy and dark matter from an unified physical origin \\cite{Kam-Mos-Pas:2001:PHYLB:}. The holographic dark energy treats vacuum energy limited by a holographic conception on degrees of freedom \\cite{Li:2004:PHYLB:}. Modified-gravity origin of the cosmic acceleration comes also from multidimensional theories of the braneworld-type cosmology \\cite{Ger:2008:PHYSR4:}. All of these models have the equation of state with the parameter $w$ varying during the cosmological expansion. Another possibility is the evolving vacuum energy density (cosmological constant) introduced in the framework of superstring quantum gravity \\cite{Ell-Mav-Nan:2000:GENRG1:}. On the other hand, the present accelerated state of the Hubble flow can be explained in the framework of the Tolman-Bondi models \\cite{Wil:2007:PHYRL:} as a consequence of inhomogeneous distribution of matter in space.  Recent observational data indicate that the allowed equation of state for the dark energy is very close to the case of a repulsive cosmological constant $\\Lambda>0$ with $\\rho_{\\Lambda}\\doteq0.745\\rho_{\\rm crit}$. Therefore, it is relevant to consider the influence of $\\Lambda>0$ in cosmological and astrophysical phenomena. The cosmological relevance of $\\Lambda>0$ is discussed in the standard textbooks \\cite{Mis-Tho-Whe:1973:Gra:,Inv:1997:IntroEinsteinRel:}.  We have studied possible effects of $\\Lambda>0$ in astrophysically motivated problems \\cite{Stu:2002:JB60:}, concentrating on radial geodesics \\cite{Stu:1983:BULAI:} and their implications in cosmological models \\cite{Stu:1984:BULAI:,Stu-Cal:1984:ASTSS1:,Stu-Sch:2006:AECIC:}, equilibrium (static) positions of spinning test particles \\cite{Stu:1999:ACTPS2:,Stu-Kov:2006:CLAQG:,Stu-Hle:2001:PHYSR4:}, and astrophysically most important disc accretion \\cite{Stu:2005:MODPLA:}, investigating circular orbits of test particles in the equatorial plane \\cite{Stu-Hle:1999:PHYSR4:,Stu-Sla:2004:PHYSR4:} and equilibrium configurations of barotropic perfect fluid tori in black-hole (BH) and naked-singularity (NS) Schwarzschild--de Sitter (SdS) and Kerr--de Sitter (KdS) spacetimes \\cite{Stu-Sla-Hle:2000:ASTRA:,Sla-Stu:2005:CLAQG:}. Elliptic integrals of  more general motion of test particles are described in \\cite{Kra:2004:CLAQG:,Kra:2005:CLAQG:,Kra:2007:CLAQG:}. The velocity profiles of the geodetical discs and perfect fluid tori in the locally non-rotating frames were studied in \\cite{Asch-etal:2004:ASTRA:,Asch:2008:CHIJAA:,Stu-Sla-Tor-Abr:2005:PHYSR4:,Stu-Sla-Tor:2007a:ASTRA:,Mul-Asch:2007:CLAQG:,Sla-Stu:2007a:CLAQG:}. Further, the test particle and perfect fluid properties were treated also in the framework of the optical reference geometry \\cite{Kov-Stu:2006:INTJM:,Kov-Stu:2007:CLAQG:,Stu:1990:BULAI:} allowing introduction of inertial forces in the intuitive Newtonian way \\cite{Abr-Car-Las:1988:GENRG2:}. Motion of photons in the Reissner--Nordstr\\\"{o}m--de Sitter and more general Kerr--Newman--(anti-)de Sitter spacetimes is discussed in the papers \\cite{Stu-Cal:1991:GENRG2:,Stu-Hle:2000:CLAQG:}. The optical phenomena were treated, e.g., in \\cite{Ser:2008:PHYSR4:,Ser:2009:PHYSR4:,Mul:2008:GENRG2:,Sch-Zai:2008:ASTRA:,Khr-Pom:2008:INTJMD:,Stu-Pls:2004:RAGtime4and5:,Bak-etal:2007:CEJP:,Stu-Hle:2002:ACTPS2:,Vir:2009:PHYSR4:}.  Self-gravitating fluid solutions of the Tolman-Oppenheimer-Volkoff equations with a non-zero cosmological constant ($\\Lambda$TOV) were first studied in \\cite{Stu:2000:ACTPS2:} for spherically symmetric stellar-like configurations with uniform energy-density distribution and generalized for cosmological solutions in \\cite{Boh:2004:GENRG2:}. Solutions of the $\\Lambda$TOV equation for polytropic and adiabatic self-gravitating configurations were investigated in \\cite{Hle-Stu-Mra:2004:RAGtime4and5:}. Stability of such stellar-like solutions was discussed in \\cite{Stu-Hle:2005:RAGtime6and7:} and \\cite{Boh-Har:2005:PHYSR4:}. The self-gravitating spherically symmetric polytropic configurations with the dark energy in the form of quintessence give interesting variety of spacetime singularities \\cite{Noj-Odi:2009:PHYLB:} that enlarges the singularities occuring in the models based on the $\\Lambda$TOV equation. Here we restrict our attention to the test fluid configurations representing physical properties of black hole backgrounds with the cosmological constant. To some extend our results could reflect also the properties of black hole solutions of the $f(R)$ gravity (for review see, e.g., \\cite{Noj-Odi:2008::}) where the cosmological constant term is present. Those solutions generally include other parameters whose influence is worth to study.  The influence of $\\Lambda>0$ on the BH-spacetime structure can be properly represented by the dimensionless cosmological parameter $y=\\Lambda M^2/3$.\\footnote{We use the geometric system of units $(c=G=1)$ hereafter.} For SdS black holes admitting existence of stable circular geodesics, which are necessary for the existence of accretion discs, the cosmological parameter $y<y_{\\rm ms,e}\\doteq 0.000237$ \\cite{Stu-Hle:1999:PHYSR4:}. Cosmological tests using the supernova magnitude-redshift relation and the measurements of CMB fluctuations \\cite{Spe-etal:2003:ASTJS:,Spe-etal:2007:ASTJS:} imply $\\Lambda_0\\sim 10^{-56}$cm$^{-2}$, and thus very low present values of $y$ for astrophysically realistic black holes. The most massive currently observed black hole seems to be harboured in the quasar TON~618, having the mass of $6.6\\times 10^{10}$\\,M$_{\\odot}$ \\cite{She-etal:2004:ASTRJ2:,Zio:2008:CHIJAA:}. In this case, $y\\doteq 4.1\\times 10^{-25}$. For primordial black holes in the very early universe, with expected high values of the effective cosmological constant, the values of $y$ can be much closer to $y_{\\rm ms,e}$. Considering the electroweak phase transition at $T_{\\rm ew}\\sim 100$\\,GeV, we obtain an estimate of the primordial effective cosmological constant $\\Lambda_{\\rm ew}\\doteq 0.028$\\,cm$^{-2}$, while at the level of the quark confinement at $T_{\\rm qc}\\sim 1$\\,GeV we obtain $\\Lambda_{\\rm qc}\\sim 2.8\\times10^{-10}$\\,cm$^{-2}$ and consequently higher values of $y$ \\cite{Stu:2005:MODPLA:,Stu-Sla-Hle:2000:ASTRA:}.  Geometrically thin (Keplerian) discs are characterized by the quasi-circular motion of test particles along stable circular orbits (geodesics) in the equatorial plane of a~given spacetime.  The inner edge of the Keplerian disc thus corresponds to the innermost stable circular orbit. In the SdS and KdS backgrounds, there is also the outer marginally stable circular geodesic which puts a~natural limit on the maximal extension of Keplerian discs in the spacetimes with $\\Lambda>0$ \\cite{Stu:2005:MODPLA:}.  Equilibrium configurations of barotropic perfect fluid tori are described by equipotential surfaces of the `gravito-centrifugal' potential $W$, which correspond to the surfaces of constant pressure in the fluid. Analyzing properties of the equipotential surfaces in the case of marginally stable tori, i.e. tori with uniform distribution of the specific angular momentum $\\ell(r,\\,\\theta)=\\mbox{const}$, orbiting in the SdS and KdS backgrounds, three astrophysically relevant phenomena were found \\cite{Stu-Sla-Hle:2000:ASTRA:,Sla-Stu:2005:CLAQG:}: (i) the outer cusp in the structure of closed equipotential surfaces, corresponding to the outermost edge of the torus, (ii) strong collimation of open equipotential surfaces near the axis of rotation of the disc that could have some relevance for the collimation of huge jets outside some giant active galaxies, (iii) for current value of the effective cosmological constant $\\Lambda_0\\sim 10^{-56}$\\,cm$^{-2}$, there is a~dimensional coincidence between extension of large galaxies and maximal extension of perfect fluid tori orbiting supermassive black holes with mass in the range of ($10^{6}$--$10^{10}$)\\,M$_{\\odot}$. For this interval of masses and current cosmological constant $\\Lambda_0$, the cosmological parameter $y$ is in the range ($10^{-34}$--$10^{-26}$). Important study of stability of fluid tori in the SdS spacetimes against the so-called runaway instability was realized in \\cite{Rez-Zan-Fon:2003:ASTRA:}, suggesting a strong stabilizing effect of the outflow of matter through the outer cusp of toroidal thick discs. Therefore, the cosmological constant strongly influences the structure of disc configurations (both the geometrically thin and thick) introducing quite naturally the outer edge of accretion discs.  It should be stressed that all of the relevant effects of $\\Lambda>0$ on the BH physics are well expressed in the SdS spacetime, since the effects of dragging of inertial frames caused by the BH rotation are concentrated in the vicinity of the BH horizon, where the influence of $\\Lambda>0$ can be abandoned for realistic values of the BH mass and the present relict cosmological constant $\\Lambda_0$. Of course, the efficiency $\\eta$ of the accretion process is influenced by the frame dragging in the innermost part of the disc, where the KdS spacetime structure is relevant. For near-extreme KdS black holes characterized by the cosmological parameter $y\\ll 10^{-6}$, $\\eta\\doteq 0.4$, which is much higher in comparison with the SdS black holes where $\\eta\\doteq 0.059$. Therefore, in studying the large scale properties of disc structures, investigation of the SdS spacetime is quite sufficient, only the accretion efficiency has to be given by the KdS spacetime structure.  One of important problems that could be hardly solved in the framework of full general relativistic (GR) approach is the influence of $\\Lambda>0$ on the structure of self-graviting discs. For the current stellar-mass and super-massive SdS black holes, the strong gravity near the BH horizon $r_{\\rm h}\\approx 2M$ weakens with the distance from the black hole growing and, at $r\\gg M$, the structure of the disc can be well described by the Newtonian theory. However, the Newtonian theory starts to lose its validity behind the so-called static radius $r_{\\rm s}=y^{-1/3}M$, where the repulsive effect of the cosmological constant starts to be relevant up to the other strong gravity region near the cosmological horizon $r_{\\rm c}\\approx y^{-1/2}M$. We expect that the pseudo-Newtonian (PN) approach, based on the replacing of the Newtonian gravitational potential by an appropriate PN gravitational potential in the framework of the Newtonian theory, could be succesfull in order to describe properly some GR effects. The PN approach enables to use directly the standard techniques developed in the framework of the Newtonian physics, e.g. the Newtonian discoseismology \\cite{Kat-Fuk-Min:1998:BHAccDis:}, and include the repulsive cosmological constant and some other relativistic effects into account.  These are the reasons that led us to introduce, in close analogy with the standard approach of Paczy\\'nski and Wiita \\cite{Pac-Wii:1980:ASTRA:} related to the Schwarzschild spacetime, a~PN gravitational potential for the SdS spacetime, and test its accuracy by comparing its predictions with the exact GR results \\cite{Stu-Kov:2008:INTJMD:}. We focused on the properties which play a~crucial role for a~geodesic motion and accretion processes in thin accretion discs, e.g. locations of the marginally stable and marginally bound circular geodesics. We concluded that the  PN gravitational potential introduced in \\cite{Stu-Kov:2008:INTJMD:} seems to be suitable for the thin disc description in astrophysically realistic SdS spacetimes with the cosmological parameter $y<10^{-25}$ \\cite{Stu-Kov:2008:INTJMD:}.  In the present paper, we study the PN description of perfect fluid tori orbiting the SdS black holes that approximate realistic thick accretion discs, and compare the PN results with the exact GR results. The influence of the discs on the spacetime geometry is assumed to be negligible. We study both the global geometrical structure of toroidal configurations and internal physical properties; the latter under simplified assumption of adiabatic tori. In section \\ref{s2}, we introduce the PN gravitational potential for the SdS spacetime and briefly summarize its significance for geodesical motion. In section \\ref{s3}, equilibrium configurations of barotropic perfect fluid tori are calculated using both the exact GR and approximate PN approaches. Namely, we  calculate loci of cusps of the equilibrium configurations, shapes of the barotropic tori and depths of the related potential wells; the GR and PN results are compared. In section \\ref{s4}, we study properties of the adiabatic tori in both the approaches, i.e., we determine mass-density, pressure and temperature profiles. In section \\ref{s5}, we calculate total mass of adiabatic tori in both the approaches and give some limits on the test-disc approximation used in our calculations. In section \\ref{dis}, behaviour of thermodynamic quantities across the adiabatic tori and the definition of the PN potential are discussed. In section \\ref{con}, some concluding remarks are presented. ",
        "conclusions": "\\label{con} It was shown that the PN gravitational potential $\\psi$ defined for SdS spacetimes in close analogy with the well known Pacz\\'{n}ski-Wiita potential \\cite{Pac-Wii:1980:ASTRA:} reflects existence of the static radius, diverges at both the black-hole and cosmological horizons, and predicts locations of both the inner and outer marginally stable and marginally bound circular orbits at the same radii as those following from the full GR theory \\cite{Stu-Kov:2008:INTJMD:}. The energy difference between the inner and outer marginally stable circular orbit, which plays a crucial role in the theory of thin discs, is very close to the relativistic result (for more details and comparison to the geodesic motion see \\cite{Stu-Kov:2008:INTJMD:}). Here, we have demonstrated that the PN potential $\\psi$ can be well applied even in the case of thick discs orbiting SdS black holes. It provides exact determination of the equipressure (equipotential) surfaces governing the shape of  toroidal discs in equilibrium configuration; to be precise, the shape of equipotential surfaces of the $\\ell = \\mbox{const}$ tori are given by identical formulae when spherical coordinates are used in the PN approach and standard Schwarzschild coordinates in the GR approach. In a given SdS background, the marginally stable toroidal configuration with $\\ell=\\ell_{\\rm mb}=\\mbox{const}$ is the most extended one. For $y\\lesssim 10^{-10}$ extension of marginally bound tori, expressed in terms of the gravitational radius $r_{\\rm h}\\sim M$, falls with mass parameter $M$  according to the law $R_{\\rm out}=r_{\\rm mb,o}\\sim M^{-2/3}$. For $\\ell=\\mbox{const}\\lesssim \\ell_{\\rm mb}$, extension of the toroidal accreting configurations fastly falls even with extremely small decrease of $\\ell$, similarly as the extension of excretion tori for extremely small increase of $\\ell$ above $\\ell_{\\rm mb}$.  We have shown that in both the PN and GR approaches there are very small differences when calculating potential barriers between the centre of the torus with $\\ell=\\ell_{\\rm mb}=\\mbox{const}$ and both its cusps. This indicates that physical properties of tori, predicted in both GR and PN approaches, could be very similar.  We have tested physical properties of $\\ell=\\mbox{const}$ (marginally stable) toroidal discs in the case of adiabatic tori with fluid equation of state corresponding to the ideal gas. We have shown that the profiles of thermodynamic quantities (mass-density, temperature, pressure) are given by the profiles of  equipotential surfaces, and are almost identical in both the PN and GR approaches. Some differences appear in the central part of the tori, where temperature, mass-density and pressure given by the PN model are from a few percent to tens percent higher as compared with the GR model for current  astrophysically relevant SdS spacetimes with $y<10^{-25}$. The central temperature is given by the potential depth $\\Delta W = W_{\\rm in}- W_{\\rm c}$ only. On the other hand, the central mass-density (and pressure) is related the total mass of the disc. In the most extended discs with $\\ell=\\ell_{\\rm mb}$, the PN and GR total masses of the disc are nearly equal, when a fixed central density is assumed, being a few percent higher in the GR model. Further, the non-relativistic fluid approximation can be conveniently  applied for description of marginally bound adiabatic tori, since the total energy density exceeds the rest energy density in restricted central parts of the tori for a few percent only.  Note that for $y\\rightarrow 0$, extension of the $\\ell=\\ell_{\\rm mb}=\\mbox{const}$ tori diverges and accretion tori with $\\ell<\\ell_{\\rm mb}$ are relevant only. Precision of the PN model with $\\ell=\\mbox{const}<\\ell_{\\rm mb}$ grows with $\\ell \\rightarrow \\ell_{\\rm mb}$. Therefore the results of the PN approach in Schwarzschild spacetimes \\cite{Muk:2002:ASTRJ2:} can be applied for adiabatic tori.  Extension of accretion (excretion) tori rapidly shrinks to the regions centered around the marginally  stable circular geodesics $r_{\\rm ms,i}$ ($r_{\\rm ms,o}$) when the torus parameter $\\ell=\\mbox{const}$ descends (rises) from the critical value $\\ell=\\ell_{\\rm mb}$.  In the moderate cases of accretion (excretion) tori, around currently astrophysically relevant SdS black holes,  defined by the condition $\\ell=\\mbox{const}=1/2(\\ell_{\\rm ms,i}+\\ell_{\\rm mb})$ ($\\ell=\\mbox{const}=1/2(\\ell_{\\rm mb}+\\ell_{\\rm ms,o})$), with the fixed limiting mass $m=M/10$, significant differences appear. We find the central mass-density and pressure to be highly dependent on the black hole mass $M$ and adiabatic index $n$ for accretion tori; the differences of the central density and pressure obtained by the GR and PN models represent tens of percent, i.e., substantially more in comparison with the most extended tori. On the other hand, their central temperature is almost independent of the black-hole mass $M$. For excretion tori, the central mass-density is almost independent of the mass of supermassive black holes ($M \\geq 10^6\\,{\\rm M_{\\odot}}$), while their central pressure and central temperature depend significantly on $M$, because of dependence of the potential depth on $M$ in such excretion tori. The GR and PN models give quite negligible differences for these central values of thermodynamic quantities, not overcoming $10^{-4}$ percent. This behaviour can be intuitively understood easily since the accretion tori under consideration are completely located in the strong gravity region near the black-hole horizon, while the excretion tori are in the region close to the static radius where the gravity is very weak and can be very well approximated by the PN potential. The most extended $\\ell=\\ell_{\\rm mb}$ tori have only a small part located in the strong gravity region.  We conclude that the PN modelling of tori in SdS spacetimes works very well in current astrophysically relevant situations with black holes of mass $M\\lesssim 10^{10}$ M$_{\\rm \\odot}$ (corresponding to the cosmological parameter $y\\lesssim 10^{-26}$) for tori with $\\ell$ close to $\\ell_{\\rm mb}$ and for all excretion discs with $\\ell_{\\rm mb}<\\ell<\\ell_{\\rm ms,o}$, since physical properties of the fluid are properly reflected in these situations by the PN approximation. The PN approach can be useful even for the accretion tori in strong gravity regions, but the relativistic phenomena can shift the GR results for thermodynamic quantities by tens percent of PN results.  We believe that the PN approach is useful in describing a wide range of astrophysical phenomena influenced by the observed relic cosmological constant. Especially, we can assume that the PN potential could be useful in developing models of excretion discs that could (potentially) serve as seeds of dark matter halos."
    },
    "0910/0910.5293.txt": {
        "abstract": "We present results of a series of Monte Carlo simulations investigating the imprint of a central intermediate-mass black hole (IMBH) on the structure of a globular cluster. We investigate the three-dimensional and projected density profiles, and stellar disruption rates for idealized as well as realistic cluster models, taking into account a stellar mass spectrum and stellar evolution, and allowing for a larger, more realistic, number of stars than was previously possible with direct N-body methods. We compare our results to other N-body and Fokker-Planck simulations published previously. We find, in general, very good agreement for the overall cluster structure and dynamical evolution between direct $N$-body simulations and our Monte Carlo simulations. Significant differences exist in the number of stars that are tidally disrupted by the IMBH, which is most likely an effect of the wandering motion of the IMBH, not included in the Monte Carlo scheme. These differences, however, are negligible for the final IMBH masses in realistic cluster models as the disruption rates are generally much lower than for single-mass clusters. As a test   to observations we model the cluster NGC 5694, which is known to possess a central surface brightness cusp consistent with the presence of an IMBH. We find that not only the inner slope but also the outer part of the surface brightness profile agrees with corresponding observations. This  provides further evidence for the possible existence of a central IMBH in NGC5694. ",
        "introduction": "As recently as 10 years ago, it was generally believed that black holes (BHs) occur in two broad mass ranges: stellar ($M_{BH}\\backsimeq3-20M_{\\odot}$), produced by the core collapse of massive stars, and supermassive ($M_{BH}\\sim10^{6}-10^{10}M_{\\odot}$), believed to have formed in the centers of galaxies at high redshift and grown in mass as the result of galaxy mergers (see, e.g., Volonteri, Haardt \\& Madau 2003). However, the existence of BHs with masses intermediate between these two ranges could not be established by observations up until recently, although intermediate mass BHs (IMBHs) were discussed by theorists more than 30 years ago; (see, e.g., Wyller 1970). Indirect evidence for IMBHs has accumulated over time from observations of so-called ultraluminous X-ray sources (ULXs), objects with fluxes that exceed the angle-averaged flux of a stellar mass BH accreting at the Eddington limit. An interesting result from observations of ULXs is that many, if not most, of them are associated with star clusters. It has long been speculated (e.g., Frank \\& Rees 1976) that the centers of globular clusters (GCs) may harbor BHs with masses $\\sim10^{3}{\\rm {M_{\\odot}}}$. If so, these BHs affect the distribution function of the stars, producing velocity and density cusps (Bahcall \\& Wolf 1976). While the detection of ULXs can only give indirect evidence of the presence of IMBHs, observations of cuspy velocity profiles would make it possible to directly determine the BH mass. However, the radius of influence of an IMBH, defined as the radius where the orbital velocity around the BH equals the velocity dispersion of the cluster, is very small. For example, at a distance of $10\\,{\\rm kpc}$, a $10^{3}{\\rm M}_{\\odot}$ BH would influence orbits within $\\approx1\"$ , making observations very challenging.  Studies of the surface density profile of GCs offer a complementary method of constraining the effects of an IMBH on the host GC stars. A recent study by Noyola \\& Gebhardt (2006) obtained central surface brightness profiles for 38 Galactic GCs from HST WFPC2 images. They showed that half of the GCs in their sample have slopes for the inner surface brightness profiles that are inconsistent with simple isothermal cores, which may be indicative of an IMBH. However, it is challenging to explain the full range of slopes with current models. While analytical models can only explain the steepest slopes in their sample, recent N-body models of GCs containing IMBHs (Baumgardt et al. 2005), might explain some of the intermediate surface brightness slopes.  However, the disadvantage of current N-body simulations is that for realistic cluster models, which take into account stellar evolution and a realistic mass spectrum, the number of stars is restricted to typically less than $\\sim10^{5}$ as these simulations require a large amount of computing time. However, many GCs are known to be very massive, with masses reaching up to $2\\times10^{6}\\,{\\rm {M_{\\odot}}}$ resulting in a much larger number of stars one has to deal with when modeling these objects. In previous N-body simulations, such large-N clusters have been scaled down to low-N systems. Scaling down can be achieved in two ways (e.g. Baumgardt et al. 2005): either the mass of the central IMBH $M_{BH}$ is kept constant and $N$ is decreased, effectively decreasing the total cluster mass $M_{C}$, or the ratio $M_{BH}/M_{C}$ is kept constant, while lowering both $M_{BH}$ and $M_{C}$. As both $M_{BH}/M_{C}$ and the ratio of $M_{BH}$ to stellar mass are important parameters that influence the structure and dynamics of a cluster, but cannot be held constant simultaneously when lowering $N$, it is clear that only with the real $N$ a fully self-consistent simulation can be achieved. Scaling becomes even more difficult once other physical processes are included into the simulations such as stellar evolution or stellar collisions. Using the correct number of stars in a dynamical simulation ensures that the relative rates of different dynamical processes (which all scale differently with $N$) are correct.  It is clear that in order to study the evolution of old globular clusters that might harbor IMBHs, more approximatee methods have to be employed. They fall roughly into two categories, methods that treat the cluster as a continuum distribution function \\citep{2004MNRAS.352..655A,1997PASJ...49..547T,1994MNRAS.269..241G,1991ApJ...370...60M,1978ApJ...226.1087C,1977ApJ...211..244L,1977ApJ...216..883B}, and Monte Carlo (MC) methods that use a particle based approach (see, e.g., \\citet{2007ApJ...658.1047F,2002A&A...394..345F,2001A&A...375..711F,1978ApJ...225..603S,1971Ap&SS..14..151H} and references therein). While the former methods have been successfully used to study the effect of a massive BH on the cluster structure \\citep{2004MNRAS.352..655A,1991ApJ...370...60M,1978ApJ...226.1087C,1977ApJ...211..244L}, the model clusters were highly idealized, consisting only of equal-mass stars, and did not incorporate stellar evolution. Although including additional physical processes is not impossible, it remains nevertheless highly non-trivial for these methods. The MC method, on the other hand, relies on a star-by-star description of the  cluster and has, therefore, the great advantage that additional processes are easily incorporated.  Our group has been developing over many years a state-of-the-art MC code, which treats many relevant processes in sufficient detail, making direct comparison with GC observations feasible (\\citealt{2007ApJ...658.1047F} and references therein). This paper is the sixth paper in a series studying the fundamental aspects of cluster dynamics using this code. Here we will describe the changes we made to our code in order to incorporate the effect of a massive central IMBH and carry out comparison runs with idealized models as well as more realistic cluster simulations published previously in the literature. In \\S \\ref{sec:Monte-Carlo-Method} we briefly describe our method and the changes we made to the MC code. We validate our code by comparing our results to previously published results in the literature, using idealized cluster models (\\S\\ref{sec:Idealized-Models}), as well as more realistic ones that include stellar evolution (\\S\\ref{sec:Realistic-Cluster-Models}). In \\S\\ref{sub:Comparison-to-Real} we present surface brightness profiles from our large-$N$ runs and compare them to observations of \\citet{2006AJ....132..447N}. We conclude in \\S\\ref{sec:Summary-and-Conclusions}. ",
        "conclusions": "}  In this paper we studied the influence of IMBHs on the evolution of globular clusters using our Monte Carlo code, which has been extended to include the effects of a central IMBH on the stellar distribution. The IMBH is treated as a fixed point mass in the cluster center, and at each time step stars are tested for entry into their loss-cone. In order to test our implementation we carried out a large number of idealized, as well as more realistic dynamical simulations of globular clusters and compared our results with published results from direct $N$-body and Fokker-Planck calculations.  In general we found that our results agree reasonably well with results from $N$-body and Fokker-Planck type codes. Significant differences only exist for idealized models while for more realistic clusters these differences become negligible.  We found that our code is able to reproduce the expected density cusp and Keplerian velocity profile around the IMBH for a single-mass cluster very well. In addition, we also found that the evolution of the cluster density profile outside of the $1\\%$ Lagrange radius is in very good agreement with direct $N$-body simulation. The only notable difference is that the cusp extends further down towards the IMBH than in the direct $N$-body simulations of \\citet{2004ApJ...613.1133B}. The reason is most likely  that the IMBH is fixed at the cluster center while in the $N$-body simulations it is allowed to move freely. Although this has a minor effect on the density profile, causing it mainly to be flatter inside $r_{w}$, we argue that it produces significant differences in the disruption rates for single-mass clusters. In particular, as we demonstrated by simple estimates of the disruption rate, the effect of IMBH wandering might be able to explain why we see  up to a factor of $2$ larger disruption rates for low-mass IMBHs than for higher-mass ones, compared to $N$-body results, which is difficult to explain otherwise. For a sufficiently fast moving IMBH the loss-cone can always be assumed full in which case the disruption rate is a strong function of the density at the disruption radius $r_{t}$. Therefore, for larger $r_{w}$, and consequently for lower $M_{BH}$, the density at $r_{t}$ becomes lower inside the very steep power-law cusp which, in turn, decreases the disruption rate. In our calculations the disruption rate is determined by the diffusion of stars into the loss-cone. This rate turns out to be larger than for the case of a wandering IMBH for our cluster parameters if $r_{w}$ is significantly larger than $r_{crit}$, which agrees with the differences we see for run BME-16. For lower $r_{w}$, and, thus, larger $M_{BH}$, loss-cone depletion effects appear to become important as we find from our simulations that the rates for these two cases converge for IMBHs with $M_{BH}/M_{*}\\gtrsim1000$, while equation (\\ref{eq:ratio-ndot-lc-w}) predicts that the rate for a wandering IMBH should become larger than for a fixed one.  Comparing the IMBH mass growth in the empty-loss-cone regime  with different Fokker-Planck type codes, such as the gas code by \\citet{2004MNRAS.352..655A} and the Monte Carlo code by \\citet{2002A&A...394..345F}, we found better agreement with our results. For all codes the IMBH growth shows the same qualitative behavior, that is, the the masses converge to a constant value. Only at late times the masses differ significantly by factors of 1.3 to 2. The larger IMBH masses in our simulations compared to the results of \\citet{2002A&A...394..345F} are most likely due to the slower expansion of the inner regions in our simulations, which are caused by the larger timesteps relative to the local relaxation time, a consequence of the shared timestep scheme used in our code. Contrary to the effect of IMBH wandering, this effect does not introduce a dependence on $M_{BH}$ and our final masses always differ by $30\\%$ regardless of the initial $M_{BH}$.  The situation for realistic cluster models is, however, rather different. In this case the disruption rates are generally lower because the cluster core gets quickly dominated by dark remnants, which have extremely small $r_{t}$. Therefore, the influence of the IMBH wandering is, in general, much reduced and does not produce significantly different results compared to a fixed IMBH. There are only minor deviations in the final cluster mass and size, with our clusters being more compact by $\\approx10\\%$ and slightly more massive by $\\approx1\\%$. The reason here might be related to the fact that we do not model close encounters with stars in the cusp, which could be an important ejection mechanism for the stars \\citep{2004ApJ...613.1133B}. Such encounters are also responsible to efficiently remove dark remnants from the cusp \\citep{2004ApJ...613.1143B} and, thus, allow for lower-mass main-sequence stars to get close to the IMBH and disrupted. The somewhat larger number of disruptions in the $N$-body simulations is also caused by the wandering of the IMBH as it can, this way, get closer to the main-sequence stars that are, due to mass segregation, further out. Although, the total mass the IMBHs gain in the $N$-body simulations is up to 3 times larger than in our simulations, the total difference amounts to just $20\\,{\\rm M}_{\\odot}$ at most or less than $7\\%$ of the total cluster mass and is, thus, very minor, given that this difference comes from not much more than 10 disrupted main-sequence stars.  Using our Monte Carlo code we are also able to reproduce the intermediate power-law slopes in the surface brightness profile seen in all $N$-body simulations with IMBH in \\citet{2005ApJ...620..238B}. Although, there is considerable scatter in the profile which makes the determination of the individual slopes rather uncertain, we  show that they, nevertheless, are within the same region as in the $N$-body simulations. Using twice as many stars in our simulations than were used in \\citet{2005ApJ...620..238B} we have a high enough resolution to obtain a reliable fit, which turns out to be close to the average slope found by \\citet{2005ApJ...620..238B}. This, once again, confirms that, despite differences in the details of the dynamics in the innermost cusp region, the overall cluster structure and evolution is rather well reproduced.  In order to see if our cluster models can  reproduce observations, we calculated the surface brightness profile from our largest $N$ run and compared it to the profile of NGC 5694 obtained by \\citet{2006AJ....132..447N}, a cluster similarly isolated from the Galactic tidal field as our models. Here we found that the overall shape of the cluster matches the observed profile up to 2 half-mass radii very well. Outside of that region the observed profile becomes flatter which is most likely due to confusion with the background. Although, the existence of an IMBH in this particular cluster seems inconsistent with the static models of \\citet{2007MNRAS.381..103M}, with NGC5694 having a concentration that is too large for the measured inner surface brightness slope, considering possible errors in the value for the concentration quoted in the Harris catalog \\citep{1996AJ....112.1487H} and the rather large uncertainty in the inner surface brightness slope, this is not necessarily the case. The fact that our model with IMBH not only reproduces the inner slope but also matches the shape of the observed surface brightness profile of NGC 5694, up to a radius of $2\\, r_{h}$, provides a stronger indication for the existence of an IMBH with a mass of the order of $0.3\\%$ of the total cluster mass.  In summary, we find that our Monte Carlo code, which now includes the effects of a central IMBH on the stellar distribution, is able to reproduce the overall structure and evolution of single-mass as well as of more realistic, multi-mass cluster models reasonably well compared to direct $N$-body simulations. There are differences in the disruption rates and IMBH masses which are most likely related to the wandering of the IMBH, which is not included in the Monte Carlo code. However, these differences become almost negligible for realistic cluster models as the disruption rates are generally very low for these cases. In addition, since the IMBH wandering radius scales with $M_{BH}/M_{*}$  as $r_{w}\\sim\\sqrt{M_{BH}/M_{*}}$  \\citep{2002ApJ...572..371C} the influence of the IMBH motion is reduced for more massive clusters with a similar $M_{BH}/M_{c}$. One process that still plays a role for more massive clusters and which is not included in our Monte Carlo simulations are the close gravitational interactions with cusp stars which we plan to implement in the near future."
    },
    "0910/0910.0072_arXiv.txt": {
        "abstract": "Using the test-field method for nearly irrotational turbulence driven by spherical expansion waves it is shown that the turbulent magnetic diffusivity increases with magnetic Reynolds numbers. Its value levels off at several times the rms velocity of the turbulence multiplied by the typical radius of the expansion waves. This result is discussed in the context of the galactic mean-field dynamo. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.3670_arXiv.txt": {
        "abstract": "Finite source effects can be important in observations of gravitational microlensing of stars. Near caustic crossings, for example, some parts of the source star will be more highly magnified than other parts. The spectrum of the star is then no longer the same as when it is unmagnified, and measurements of the atmospheric parameters and abundances will be affected. The accuracy of abundances measured from spectra taken during microlensing events has become important recently because of the use of highly magnified dwarf stars to probe abundance ratios and the abundance distribution in the Galactic bulge. While the abundance ratios in general agree with the giants, with the possible exception of [Na/Fe], there may be more dwarfs than giants with high [Fe/H]. In this paper, we investigate the effect of finite source effects on spectra by using magnification profiles motivated by two events to synthesize spectra for dwarfs between 5000K to 6200K at solar metallicity.  We adopt the usual techniques for analyzing the microlensed dwarfs, namely, spectroscopic determination of temperature, gravity, and microturbulent velocity, relying on equivalent widths. We find that ignoring the finite source effects for the more extreme case results in errors in \\teff $< 45$K, in \\logg{} of $<$0.1 dex and in $\\xi$ of $<$0.1 km/s. In total, changes in equivalent widths lead to small changes in atmospheric parameters and changes in abundances of $<$0.06 dex, with changes in [\\ion{Fe}{1}/H] of $<$0.03 dex. For the case with a larger source-lens separation, the error in [\\ion{Fe}{1}/H] is $< 0.01$ dex. This latter case represents the maximum effect seen in events whose lightcurves are consistent with a point-source lens, which includes the majority of microlensed bulge dwarfs published so far. ",
        "introduction": "}  A microlensing event occurs when a stellar mass object passes between the observer and a background source. For events in the Local Group, the source is a star. If the source is finite, rather than infinitely small, parts of the source may be magnified more than other parts. In the case of stars, for example, the limb can be brightened relative to an unmagnified source. Since the light from the limb comes from different temperature profiles in the star compared to the center, this will affect the spectrum \\citep[e.g.][] {valls-gabaud:95,loeb:95} For example, a spectrum of the limb of the Sun has Balmer lines with weaker wings, while lines of other elements can be either weakened or strengthened relative to the center \\citep{hastings:73,hale:07}.  Recently, observers using large telescopes have published abundance ratios in dwarfs in the Galactic bulge from spectra obtained while the sources were highly magnified.  A number of these have turned out to have super-solar metallicities, including OGLE-2006-BLG-265S \\citep{johnson:07}, MOA-2006-0BLG-099S \\citep{johnson:08}, OGLE-2007-BLG-349S \\citep{cohen:08}, MOA-2008-BLG-310S and MOA-2008-BLG-311S \\citep{cohen:09} and OGLE-2007-BLG-514S \\citep{epstein:09}. Some are metal-poor including the subgiant OGLE-2008-BLG-209S \\citep{bensby:09a} and the dwarf OGLE-2009-BLG-076S \\citep{bensby:09b}, which, at [Fe/H]$=-0.76$ is the most metal-poor dwarf/subgiant observed when microlensed. For a sample of eight microlensed dwarfs, \\citet{epstein:09} found that a K-S test gave a 1.6\\% chance that the dwarfs were drawn from the same metallicity distribution function (MDF) as the giant MDF from \\citet{zoccali:08}. \\citet{bensby:09c} report results for a total of 13 dwarfs, including new results for more metal-poor dwarfs, so clearly the comparison between the giants and the dwarfs is an evolving topic. \\citet{cohen:08} proposed that mass-loss on the giant branch prevents some more metal-rich stars from becoming red giants, similar to the mechanism suggested by \\citet{kalirai:07} to explain the low-mass He white dwarfs in the metal-rich cluster NGC 6791. However, this explanation predicts that the red clump MDF is biased to more metal-poor stars because of metal-rich stars skipping the horizontal branch phase and that the red giant luminosity function should drop more than expected from theoretical models, predictions that do not agree with the giant data \\citep{zoccali:08}. Another explanation is that there are systematic differences between giant and dwarf abundance analyses.  \\citet{cohen:09} calculated that a systematic offset of 0.10 dex between the dwarf and giant metallicity scales would give the mean metallicity offset a significance of 2-$\\sigma$ and larger systematic offsets would obviously decrease the significance further still. Finally, \\citet{zoccali:08} suggested that the spectra of microlensed dwarfs could be affected by differential magnification sufficiently that the usual analysis of the dwarfs, which does not take this into account, could lead to biased answers and help explain the possible discrepancy.  This last suggestion can be tested by comparing the answers obtained from synthetic spectra with and without differential limb magnification.    While most microlensing events follow the lightcurve of a point source, about $\\sim$3\\% \\citep{witt:95} of events show finite source effects. This fraction is even higher for high-magnification events that are targets of the current generation of dwarf studies. Of the eight published events for which dwarf spectra were obtained, we know that at least three of the events were affected by finite source effects, namely OGLE-2007-BLG-514, OGLE-2007-BLG-349, and MOA-2008-BLG-310. We need to determine the size of the effect that differential limb magnification (DLM) has on the spectra and the measurement of the effective temperature (\\teff), gravity (\\logg), metallicity ([Fe/H]), microturbulent velocity ($\\xi$) and abundance ratios to correctly interpret these events.  In addition to probing the chemical evolution of the bulge, the accuracy of the measured \\teff{} from the spectrum is important for testing the method by which colors of source stars in microlensing events are determined. Using a metallicity and a \\teff{} derived from the spectrum, we can predict the color of the star using relations between color and \\teff, such as that by \\citet{ramirez:05}, and compare with the color estimated using standard microlensing techniques, which rely on the offset of the star from the red clump. The results so far indicate that if the color of the red clump is $(V-I)_0$=1.05, the \\teff s derived spectra are in agreement with \\teff s from colors.  Much work has been done on the effects of DLM of the disk of giants during a microlensing event because these events are easier to find and are longer-lasting than similar events in dwarfs. \\citet{valls-gabaud:98} predicted the effects on the spectrum of a giant, in particular the H$\\beta$ and CO lines, by fitting simple analytic expressions to the limb profiles for both continuum and lines from the \\citet{kurucz:92} models. He found that the equivalent widths could change by 20\\% over the course of the event. \\citet{heyrovsky:00} computed more realistic models avoiding the use of linear approximations. They calculated contribution functions for lines, which indicate where the lines are formed, and illustrated the range of behavior individual spectral features will have because of varying center-to-limb profiles. \\citet{gaudi:99} pointed out that binary lens events had particular advantages for resolving stellar surfaces, for example the fact that the second caustic crossing can be predicted if the event is monitored, and therefore carefully observed.   Thanks to intensive monitoring by observers, the effect of the size of the source on microlensing events has been observed many times. Finite source effects were observed for the giant MACHO 95-30 \\citep{alcock:97}, including variations in the equivalent widths of H$\\alpha$ and the TiO bands. The bulge K3 giant star EROS BLG-2000-5S was highly magnified by a binary lens. Because the caustics of a binary lens are symmetric, the timing of the second caustic crossing was predicted to high accuracy and allowed intensive follow-up during this event. Both \\citet{albrow:01} and \\citet{castro:01} found decreases in the strength of the H$\\alpha$ line when the limb was more highly magnified relative to the center. \\citet{afonso:01} monitored the same event photometrically, measured the color of the limb, and found that the outer 4\\% of the limb showed strong H$\\alpha$ emission, which they ascribed to the presence of a chromosphere. \\citet{albrow:99} performed the first limb-darkening measurement of a bulge giant for the event MACHO 1997-BLG-28, using photometry in the $V$ and $I$ bands.  Other limb-darkening measurements were done for giants including MACHO 1997-BLG-41 \\citep{albrow:00}, and OGLE-2002-BLG-069S \\citep{kubas:05}, which the authors found agreed with a linear limb-darkening law. An opposite result was found by \\citet{cassan:06} who obtained photometry for the K giant OGLE-2004-BLG-254S and derived limb darkening coefficients in the $I$ and $R$ bands. Combining these results with other microlensed-based limb darkening measurements of K giants and comparing them to the predictions of ATLAS9 models, they found a disagreement, which they suggested was the result of the inadequacy of the linear limb darkening law. Time-resolved spectra of OGLE-2002-BLG-069S, a G5III star, by \\citet{cassan:04}, showed the power of this technique for probing the stellar atmosphere. In this case, the changes during the event in  wings of H$\\alpha$ line, which are formed deep in the atmosphere, agreed with model predictions while the core of the line showed an emission peak, which again can only be explained by a chromosphere. That the temperature structure in the outer layers is not well understood for this event was also found by the analysis of \\citet{thurl:06}. These studies show the potential to investigate the atmospheric structure of giants.  Studies of dwarf stars have been much rarer. \\citet{afonso:00} and \\citet{abe:03} reported the only measurements of limb darkening in dwarfs, for MACHO 98 SMC-1, a main-sequence A star and MOA 2002-BLG-33S, a solar-type star, respectively. The most relevant work to this paper is the analysis of \\citet{lennon:96} of the event MACHO 96-BLG-3. They took three exposures when the source, a dwarf star, was moving across a caustic. These R$\\sim$1100 spectra with signal-to-noise (S/N) of 25-100 did not show any convincing cases of profile variability. \\citet{lennon:96} compared the observations of this star with both a library of high-resolution, high S/N spectra of F and G stars observed at Calar Alto and a grid of synthetic spectra, and derived stellar parameters of \\teff=6100K from the H$\\alpha$ line, \\logg=4.25 from the \\ion{Mg}{1} triplet and a metallicity ([M/H]) between 0.3 and 0.6 from fits to regions with many \\ion{Fe}{1} and \\ion{Ca}{1} lines. They calculated the expected deviations from the unmagnified spectrum for this event, and found that the expected change in the line profiles for the three spectra was $\\leq$ 1\\%, while the expected change in the continuum was $\\leq$ 2\\%. Given the small changes expected and the S/N of the spectra, it is not surprising that no changes were observed, and that the derived atmospheric parameters of the star would also not be affected. Indeed, it is not surprising that most changes are small. The amount of magnification depends on the distance from the lens, but the emitted spectrum is the same for an entire annulus (for a spherical star). Because the distance from the lens varies around the annulus, the average magnification of a spectrum at a particular annulus is smaller than the largest magnification of an particular spot would suggest. In addition, while the spectrum of the star increasingly changes from center to limb, the intensity of the limb is lower than the center, by factors of a few. Therefore it is difficult to overcome the influence of the light coming from the central regions of the star in the disk-averaged light.   The work discussed above focused on learning about the atmosphere of the star and showed that finite source effects cause changes in the spectrum of the star, which can be inverted to give temperature as a function of depth. Because the same effects that make these extreme events interesting as probes of the atmosphere will also change the spectrum that is analyzed for elemental abundances, the goal of this paper is to determine the effect of DLM on the \\teff, $\\xi$, \\logg, [Fe/H] and abundance ratios using the standard spectroscopic techniques, such as equivalent width analysis. Therefore, we will be concerned with mimicking standard abundance analysis as closely as possible, rather than exploring the full range of knowledge that can be gleaned from time-resolved spectroscopy of microlensed events. ",
        "conclusions": "Microlensing of bulge dwarfs provides the exciting opportunity to obtain high S/N and high-resolution spectra of stars that otherwise would be unfeasible with current telescopes. However, microlensing can produce finite source effects that affect the spectra by differentially magnifying annuli in the star compared to an unmagnified star. We have examined the practical impact of such DLM by synthesizing spectra of dwarfs with magnification profiles similar to events that have been observed and performing an abundance analysis. We find that, as expected, there are changes to the spectrum, which results in changes in the atmospheric parameters, metallicities and abundance ratios. Given the small size of the changes, the S/N of the observed spectra and the accompanying error in the atmospheric parameters, and the length of the exposures compared to the duration of a specific magnification profile, the resulting effect on the abundances is small compared to other sources of error. We note that the change in abundances can be either positive or negative (see Profile 35 vs. Profile 18) depending on the form of the magnification profile. Microlensing events are more commonly in the shape of Event A, with mild magnification causing limb brightening, so there will be a tendency for events to be biased to lower metallicity,but Figure 11 shows that Event A leads to changes of $<0.01$ dex in metallicity. These results, combined with the fact that finite source effects do not affect all of the high-magnification cases for which dwarfs have been observed, demonstrate that this is not the cause of any differences between the observed metallicity distribution function for giants and dwarfs. Finally, if the lightcurve shows even more extreme finite source effects than have been modeled in this paper, magnification profiles for that event can be constructed and model spectra calculated that appropriately take those profiles into account, so this will not be a limiting problem for the accuracy of abundances derived for microlensed dwarfs."
    },
    "0910/0910.4099_arXiv.txt": {
        "abstract": "We present a study of the correlations between spectral, timing properties and mass accretion rate observed in X-rays from the Galactic Black Hole (BH) binary GRS 1915+105 during the transition between hard and soft states. We analyze all transition  episodes from this source  observed with { Rossi X-ray Timing Explorer} ({\\it RXTE}), {coordinated with Ryle Radio Telescope (RT) observations. We show that  broad-band energy spectra of GRS~1915+105 during all these spectral states can be adequately presented   by  two  Bulk Motion Comptonization (BMC) components: a  hard component  (BMC1,  photon index $\\Gamma_1=1.7-3.0$)  with turnover at  high energies and soft thermal component  (BMC2, $\\Gamma_2=2.7-4.2$) with characteristic color temperature $\\le$~1 keV, and the redskewed iron line (LAOR) component. } We also present observable correlations between the index and the normalization of the  disk ``seed'' component. The use of ``seed'' disk  normalization, which is presumably proportional to mass accretion rate in the disk, is crucial to establish the index saturation effect during the transition to the soft state. We discovered the  photon index saturation of the soft and hard spectral components at values of $\\lax$ 4.2 and 3  respectively. We   present a physical model which explains the index-seed photon normalization correlations. We argue that the index saturation effect of the hard component (BMC1)  is due to the soft photon Comptonization in the converging inflow close to BH and that of soft component is due to matter accumulation in the transition layer when mass accretion rate increases. Furthermore we demonstrate  a  strong correlation between equivalent width of the  iron line and radio flux in  GRS~1915+105. In addition to our spectral model components we  also  find  a strong feature of ``blackbody-like'' bump which  color temperature is about 4.5 keV  in  eight  observations of the intermediate and soft states. We discuss  a possible origin of this   ``blackbody-like'' emission. ",
        "introduction": "The study of the characteristic changes in spectral and variability properties of X-ray binaries is proved to be a valuable source of information on the physics governing the accretion processes and on the fundamental parameters of black holes (BHs).    The simultaneous study of the spectral and timing evolution of a BH source during a state transition has been a subject of many investigations [see references in a review by  \\cite{rm}]. \\cite{fb04}, hereafter FB04,  introduced a classification  of  the spectral states in GRS 1915 +105 and studied  the spectral state evolution.   Using X-ray colors (hardness ratio) they introduced three spectral states. State A: in which the strong blackbody-like (BB)  component of color temperature $\\gax 1$ keV dominates in the overall spectrum and little time variability  is detected. State B: similar to state A but substantial red-noise variability  on scales $>1$ s occurs in this state. State C: the spectra are harder than those in states A and B. Photon indices of the power-law components vary from 1.8 to 2.5.  White-red noise (WRN) variability  on scales $>1$ s  takes place in this state.  Furthermore FB04 discussed the connection between states A, B, C observed in GRS 1915 with the three ``canonical'' states in black hole candidates (BHCs) also identified by their timing and spectral properties.  At a low luminosity state the energy spectrum is dominated by a hard Comptonization component combined (convolved) with  a  weak thermal component. The  spectrum of this low (luminosity) hard state (LHS) is presumably a result of  Comptonization (upscattering) of soft  photons, originated  in a relatively weak accretion disk,  off electrons of  the hot ambient plasma [see e.g. \\cite{st80}]. Variability in LHS is high (fractional root-mean-square variability is up to 40\\%) and  presented by a flat-top broken power law (white-red noise) shape, accompanied by quasi-periodic oscillations (QPOs) in the range of 0.01-30 Hz, observed as narrow peaks in the power  density spectrum (PDS). In high soft state  (HSS) a photon spectrum is characterized by a prominent thermal component which is probably a signature of  a strong emission coming from a geometrically thin accretion disk. A weak power-law component is also present at the level of not more than 20\\% of the total source flux. In the HSS the flat-top variability ceases, QPOs disappear and PDS acquires a pure power-law shape. The total  variability in HSS is usually about 5\\% fractional rms. The intermediate state (IS) is a transitional  state between LHS and HSS. Note in addition to LHS, IS and HSS sometimes very soft state (VSS) is observed in which the BB  component is dominant and the power-law component is either very  weak or  absent at all. The  bolometric luminosity in VSS is a factor of 2-3 lower than that in HSS.  FB04 concluded  that probably all three states  A, B, C of GRS 1915+105 are instances of something similar to the HSS/IS observed in other BHC systems, associated to the high accretion rate value for this source, although during the hardest intervals LHS might be sometimes reached. We come to the similar conclusions analyzing spectral and timing  data from GRS 1915+105 obtained by {\\it RXTE} (see below).  Close correlations of nearly periodic variability [quasi-periodic oscillations (QPO)] observed during  low-hard and intermediate states  with the photon index of the Comptonization spectral component have been reported in multiple state transitions observed from accreting BHs [see \\cite{vig}, \\cite{ST06}, (2007),  (2009), hereafter V03,  ST06, ST07 ST09 respectively]. The ubiquitous nature of  these correlations suggests that the underlying physical processes which lead to the observed variability properties are closely tied to the corona; furthermore, they vary in a well defined manner as the source makes a transition between spectral  states. The fact that the same correlations are seen in many sources, which vary widely in both luminosity (presumably with mass accretion rate) and state, suggests that the physical conditions controlling the index and the low frequency QPOs are characteristics of these sources. Moreover,  they  may be an universal property  of  all  accreting compact systems, including neutron sources too [see \\cite{tf04}, hereafter TF04, and \\cite{ts05}].    When a BH is in LHS, radio emission is also detected and a jet is either seen or inferred \\citep{f01}. Several models are successful in reproducing the energy spectrum from the radio domain to the hard X-rays [see e.g. \\cite{MFF01}, \\cite{vrc01}, \\cite{cf02}, \\cite{M03} and \\cite{g05}]. The multiplicity of models that can fit well the time average spectrum of galactic BHs indicates that this alone is not enough to distinguish the most realistic one among them. X-ray timing features can be  the key features to finally finding the common physical connection between the  corona, the accretion disk and the jet radio emission in BHs.  There is a big debate in the literature on the origin of quasiperiodic oscillation (QPO) frequencies [see e.g. \\cite{rm}] and its connection with the radio emission. \\cite{mfk05}  reported on correlation between radio luminosity and X-ray timing features  in X-ray binaries containing  a number of low magnetic field  neutron stars and one black hole GX 339-4. They showed that in the low-hard state (LHS)   radio luminosity is correlated with the low frequency QPO  (LFQPO).  Note ST09 demonstrate that in  LHS of Galactic BHs    LFQPO  changes by order of magnitude, from 0.2 to 2 Hz whereas the photon index has almost the same value of 1.5.  Below we show that in GRS 1915+105 the photon index monotonically increases with LFQPO and disk mass accretion rate, although  the radio luminosity does not correlate with LFQPO and X-ray luminosity in the whole range of spectral states, from low-hard to high-soft through intermediate states. Recently \\cite{kyl08} suggested a model  which explains how  the QPO phenomenon is  related to an appearance of radio flares (jets).  Below (see \\S 3) we present details of our observational study  of the QPO connection with the X-ray and radio flaring activity in GRS 1915+105.   In LHS and IS  which we consider in our study, only a small part of the disk emission component is seen directly. The energy spectrum is dominated by a Comptonization component presented by a power law. To calculate the total normalization of the ``seed'' disk blackbody component we model the spectrum with a Generic Comptonization model [BMC XSPEC model, see details in \\cite{bmc}] which consistently convolves a disk blackbody with a Green function of the Compton Corona to produce the Comptonization component. We argue that the disk emission normalization calculated using this approach produces a more accurate correlation with respect to the correlation with the direct disk component which was obtained using the additive model, multicolor disk plus power law [see e.g. \\cite{mr}].   This Paper is a continuation of the study of index-QPO and index-seed photon normalization correlations in BH sources started in  ST07 and ST09.  Particularly here we present a study of the   index-seed photon normalization (disk flux) correlation observed from GRS 1915+105 when it evolves from LHS to HSS.  The description of  \\textit{RXTE} data-set used is given in \\S \\ref{data}.  We have analyzed a broader sample of state transitions from GRS 1915+105 and  we found a diverse phenomenology for index evolution through a transition.  In  \\S \\ref{obs_transition} we provide a detailed description of state transitions analyzed in this study. In \\S  4 we discuss and interpret the results of our observational study. Specifically in  \\S 4  we consider the effect of the bulk motion Comptonization in the inner part of the accretion flow on the index evolution during a state transition. Also we show that the index saturation effect is a direct consequence of the existence of this inner bulk motion region and, therefore, can be considered as an observational signature of the converging flow (black hole). Furthermore in  \\S 4 we  discuss  the TF04 model and  the Monte Carlo simulations  by  \\cite{lt09} (in preparation) in which the observable index evolution with  $\\dot m$  has been already  predicted. Conclusions follow in \\S \\ref{summary}. ",
        "conclusions": "} We concentrate our efforts on the study of the correlation between the spectral index and the  accretion disk luminosity. We argue that the shape of the correlation pattern can contain the direct BH signature. Namely, we show both observationally and theoretically that the index saturates with mass accretion rate  which is a signature of a converging flow. This index saturation effect  can exist only in  BH sources. Also this  correlation pattern carries the most direct information on the BH mass and the distance to the source (see ST09).  We compiled the state transition data from GRS 1915+105 collected with the \\textit{RXTE} mission. We examined the correlation between the photon index of the Comptonized spectral component, its normalization and the QPO frequency (see Figs.~\\ref{outburst_97_rise}-\\ref{outburst_05-06_decay}, \\ref{outburst_index_norm}-\\ref{outburst_index_qpo}).  The spectral data of GRS 1915+105 are well fitted by two  (soft and hard)  BMC components for most of analyzed IS and HSS spectra (see Fig. \\ref{two_BMC}) while LHS spectra essentially require only one BMC component. In addition to two BMC components  8  IS-HSS spectra require   an extra component which can be fitted by `` high temperature BB-like\"  profile.  We suggest this ``BB'' component is probably a signature of the redshifted annihilation line formed in the very narrow shell near a BH horizon due to high photon compactness  taking place during  intermediate  and high-soft states  (see Fig. \\ref{sp_bbody} and Table 7).  A remarkable result of our study is that  the index - normalization (mass accretion rate) correlation seen  in GRS 1\\-9\\-1\\-5+105 is predicted  by the theory of the converging flow. We  demonstrate that a strong index saturation vs disk flux  seen in the index-disk flux correlation   (see Fig.~\\ref{outburst_index_norm}) is an observational signature of the presence of the converging  flow, which should only exist  the BH sources. In other words, {\\it this index saturation effect provides a robust observational evidence for the presence of black hole in GRS~1915+105}.  We also  find  a tight positive correlation of QPO frequencies with the index  (see Fig. \\ref{outburst_index_qpo}) and consequently that with the disk flux.  Our comprehensive analysis of  X-ray and radio emissions in GRS 1915+105 shows that   QPOs are seen independently of radio activity of the source during the spectral transition from low-hard to high-soft state.  Specifically these QPO features have been  detected at any level of   the radio flux and even when the radio emission is at the noise level (see left and right panels in Fig.~\\ref{radio_QPO_independence} correspondingly). We also do not find any correlation between X-ray and radio fluxes  and X-ray power-law index. (see Fig. \\ref{outburst_index_radio_X-ray}). However,  we establish  a  strong correlation between equivalent width of iron line and radio flux in binary GRS~1915+105 (see Fig. \\ref{outburst_EW_radio}).    We are grateful to the referee whose constructive suggestions help us to  improve the paper quality. We  acknowledge productive discussion with Nikolai Shaposhnikov and  Ada Paizis and we also would like to thank Guy Pooley who kindly provided us  {\\it Ryle Radio Telescope} data.  \\appendix"
    },
    "0910/0910.1733_arXiv.txt": {
        "abstract": "Fully sampled degree-scale maps of the $^{13}$CO 2--1 and CO 4--3 transitions toward three members of the Lupus Molecular Cloud Complex --- Lupus~I, III, and IV --- trace the column density and temperature of the molecular gas. Comparison with IR extinction maps from the c2d project requires most of the gas to have a temperature of 8--10\\,K. Estimates of the cloud mass from $^{13}$CO emission are roughly consistent with most previous estimates, while the line widths are higher, around 2~km\\,s$^{-1}$. CO 4--3 emission is found throughout Lupus~I, indicating widespread dense gas, and toward Lupus~III and IV. Enhanced line widths at the NW end and along the edge of the B\\,228 ridge in Lupus~I, and a coherent velocity gradient across the ridge, are consistent with interaction between the molecular cloud and an expanding \\ion{H}{1} shell from the Upper-Scorpius subgroup of the Sco-Cen OB Association. Lupus~III is dominated by the effects of two HAe/Be stars, and shows no sign of external influence. Slightly warmer gas around the core of Lupus~IV and a low line width suggest heating by the Upper-Centaurus-Lupus subgroup of Sco-Cen, without the effects of an \\ion{H}{1} shell. ",
        "introduction": "The Lupus star forming region, recently reviewed by \\citet{comeron}, lies about 150~pc from the Earth \\citep{08alombardi}, in the Gould Belt. It is immediately visible by inspection of optical photographs of the sky, comprising a set of largely filamentary dark clouds. Based on optical extinction maps \\citep{cambresy}, the `Cores to Disks' Spitzer Legacy Programme \\citep[c2d,][]{c2d,c2dproducts,c2dresults}, has produced and analyzed infrared images of three areas: Lupus~I, III, and IV\\footnote{Some authors use arabic numerals: Lupus~1, 3 and 4}. The c2d data products \\citep{c2dproducts} include: shorter-wavelength Infrared Array Camera (IRAC) maps, used to identify and classify young stellar objects \\citep[YSOs,][]{merin}; far-IR MIPS maps \\citep{mips} tracing the thermal emission of the dust component of the molecular clouds; and IR extinction maps, based on the Two Micron All Sky Survey \\citep[2MASS;][]{mips} and on a combination of 2MASS and IRAC data \\citep{c2dproducts}. In this paper, we present maps of the c2d fields in Lupus, tracing the molecular hydrogen by the rotational emission of carbon monoxide (CO) and isotopically substituted \\thco.  The Lupus complex has been extensively surveyed in the various 1--0 transitions of CO: \\citet{86murphy} first mapped the Lupus molecular clouds using the 1.2~m Columbia telescope with 0.5\\degr\\ resolution to show the extent of the molecular gas in Lupus and suggest a total mass of a few $10^4$~$M_\\sun$. Improved maps in several transitions were obtained by NANTEN (2.7\\arcmin\\ resolution, but routinely undersampled): \\citet{01tachi} published \\twco\\ maps of the whole complex, mostly sampled at 8\\arcmin\\ spacing, with some areas at 4\\arcmin\\ spacing; \\thco\\ 1--0 maps \\citep{96tachi} cover Lupus~I and III with 8\\arcmin\\ spacing, while \\ceo\\ 1--0 maps \\citep{99hara} cover all the clouds in the complex at 2\\arcmin\\ spacing. \\citet{02my} mapped about 200 arcmin$^2$ of Lupus~IV at sub-arcminute resolution and full-beam spacing in the 1--0 transitions of CO, \\thco, and \\ceo. The \\thco\\ 2--1 and CO 4--3 maps presented here are the first large-scale fully sampled low-$J$ CO maps of Lupus, and the first mid-$J$ CO maps of any kind.  The large-scale structure of the Lupus clouds can be traced by near-IR extinction \\citep{08blombardi} and CO 1--0 emission \\citep{01tachi}. The complex covers some 20\\degr\\ of galactic latitude (about 50~pc). At low latitudes ($b\\la 10^\\circ$), a large mass of diffuse gas contains denser filamentary clouds; at higher latitudes, the dense clouds (Lupus I and II) are more clearly separated. The evolution of the Lupus clouds may have been driven by the influence of nearby OB stars \\citep{01tachi}. The Lupus and Ophiuchus molecular clouds face one another across the Upper-Scorpius subgroup of the Scorpius-Centaurus OB Association; the Upper-Centaurus-Lupus subgroup lies on the opposite side of the Lupus clouds from Upper-Sco. The \\ion{H}{1} shell around Upper-Sco, blown by stellar winds and a presumed supernova 1.5\\,Myr ago \\citep{92degeus}, borders the NE side of the Lupus clouds; on the plane of the sky, the ridge which dominates Lupus~I (B\\,228, see below) lies just on the trailing edge of the \\ion{H}{1} shell. The much older Upper-Cen-Lup subgroup has driven an \\ion{H}{1} shell far beyond the Lupus clouds \\citep{92degeus}; it should have passed through the Lupus complex some 4--7 Myr ago, roughly consistent with the ages of T~Tauri stars in Lupus III and IV \\citep{02my}.  Lupus~I (Figure~\\ref{fig-lui-opt}) is dominated by B\\,228 \\citep{barnard} or GF\\,19 \\citep{79schneider}, a long ridge running NW--SE, extending over about 2\\degr\\ (5~pc) parallel to the edge of the Upper-Sco \\ion{H}{1} shell. The molecular material appears to fall off steeply toward the center of the shell (at the NE edge), but there is extensive material on the other side of the ridge. Lupus~III (Figure~\\ref{fig-luiii-opt}) is a long (about $4\\times 1$~pc) E--W cloud \\citep[GF\\,21,][]{79schneider} at the edge of the low-latitude cloud mass: at its eastern end, it curves up to the NE and breaks up into clumps; toward the west lie two embedded (but optically visible) Herbig Ae/Be stars, HR\\,5999 (A\\,5--7) and the B\\,6 HR\\,6000 \\citep{comeron}. \\citet{02tachi} consider Lupus~III to be a cluster-forming structure with a dense head and spread-out tail, analogous to the $\\rho$~Oph and Cha~I clouds. A couple of pc north of Lupus~III, there is another smaller cloud, which appears to be a separate condensation in the low-latitude gas; we henceforth refer to this northern structure as Lupus~III\\,N. Lupus~IV (Figure~\\ref{fig-luiv-opt}) is the head of a filamentary structure running approximately E--W \\citep[GF\\,17,][]{79schneider,02my}, comprising a small core, about 0.5~pc long (E--W), lying within a more diffuse extended cloud, about a pc across. The filament connects Lupus~IV to the low-latitude gas mass to the East; \\citeauthor{79schneider} describe it as a ``chain of small faint globules''.  Individual dark clouds or cloud peaks have been identified by inspection of optical surveys on scales ranging from degrees (e.g., B\\,228) to arcminutes \\citep{sandy,feitzinger,86hartley,bhr1,96avb,99lm,vmf}; these clouds are listed in Table~\\ref{tab-dcs}. Near-IR extinction maps of Lupus~III with a resolution of 30\\arcsec\\ \\citep{05teix} have been used to identify dark cores in more detail, which were then fitted with Bonnor-Ebert profiles. Single-pointing spectra in several molecular lines have been observed toward extinction-selected clumps: \\citet{sandy} observed H$_2$CO along four lines of sight at 6~cm (6.6\\arcmin\\ beamwidth); \\citet{bhr2} searched for ammonia, but did not detect it, toward two globules (BHR\\,120 \\& 140); \\citet{vmf} observed \\thco\\ and \\ceo\\ 1--0, but their 0.8\\arcmin\\ beamwidth is poorly-matched to ours; \\citet{04lmp} observed a few of their Lupus~I cores \\citep{99lm} in CS 3--2 and \\dcop\\ 2--1 (0.7\\arcmin\\ beamwidth). Previously identified CO cores, both those observed by \\citeauthor{vmf}, and those identified from their \\ceo\\ maps by \\citet{99hara}, are listed in Table~\\ref{tab-mcs}. The positions of these extinction cores and \\ceo\\ cores with respect to optical and \\thco\\ 2--1 emission are shown in Figures~\\ref{fig-lui-opt}--\\ref{fig-luiv-opt}. Of the members of the c2d small dark cloud sample mapped by AST/RO in \\thco\\ 2--1 and CO 4--3 \\citep{andrea}, three lie within the Lupus region, but outside our maps.  Lupus~I contains a few YSOs, concentrated to the B\\,228 ridge \\citep{merin}. Lupus~III contains a very dense cluster of YSOs near the Herbig Ae/Be stars, composed mainly of T~Tauri stars \\citep{07allen,merin}, but including the Class~0 object Lupus\\,3\\,MMS \\citep{07tachi}. \\citet{09comeron} identified a population of cool stars and brown dwarfs toward Lupus~I and III, which are not strongly concentrated toward the molecular clouds. The Lupus~IV region contains about as many YSOs as the Lupus~I region \\citep{merin}, but they are older (Class II/III) and largely found outside the molecular cloud, with the exception of one flat-spectrum YSO candidate close to the cloud center. Even the dispersed population is absent \\citep{09comeron}. Lupus~IV thus appears to have less (or no) ongoing star formation, compared to Lupus~I and III. \\citeauthor{merin} suggest that the Class II and III YSOs surrounding it represent an earlier generation of star formation (possibly associated with the passage of the Upper-Cen-Lup \\ion{H}{1} shell mentioned above), and that the very dense Lupus~IV core is poised to form new stars. The younger YSOs \\citep[Class 0, I, and flat-spectrum sources from the c2d samples,][]{merin} within the boundaries of our \\thco\\ 2--1 maps are listed in Table~\\ref{tab-yso}, and plotted in Figures~\\ref{fig-lui-opt}--\\ref{fig-luiv-opt}. ",
        "conclusions": "Fully sampled degree-scale maps of the \\thco\\ 2--1 emission toward the Lupus I, III, and IV clouds trace the column density and temperature of the gas, the transition becoming optically thick in the cloud cores. The peak \\trstar\\ is well correlated with the near-IR extinction \\citep{mips,c2dproducts}, and a comparison of the two suggests that the bulk of the molecular gas in Lupus has a temperature of 8--10\\,K, rather than the 10--17\\,K generally adopted elsewhere \\citep[e.g.][]{96tachi,02my,05teix}. This estimate is fairly robust to changes in the relationship between \\thco\\ column density and \\av. Estimates of the cloud masses from the \\thco\\ maps are reasonably consistent with those derived from extinction mapping. The differences between these estimates vary greatly from cloud to cloud, and suggest that there may be significant spatial variation in the \\thco-to-\\av\\ relationship, as found in the Perseus complex \\citep{08goodman}. The line widths of \\thco\\ 2--1 toward the clouds are higher than previous estimates \\citep{c2dresults}, around 2\\,\\kms, with Lupus~III\\,N rather broad and Lupus~IV rather narrow.  Fully sampled CO 4--3 maps covering most of Lupus~I and small regions of Lupus~III and IV trace dense gas: the peak \\trstar\\ generally indicates excitation temperatures quite close to those of \\thco\\ 2--1, and hence that the transition is largely thermalized. This suggests that the volume density $\\ga 10^4$\\,\\percc, although modelling will be required to ascertain the required density. CO 4--3 emission is pervasive toward Lupus~I (the map of which covers a large area), implying that this dense gas is found either throughout or all around the outside of the cloud, although it may comprise a fairly small fraction of the cloud mass.  The physical conditions of the molecular gas vary along the B\\,228 ridge in Lupus~I. At the NW end, the gas has broader lines and probably higher temperature than in the bulk of the cloud; the column density is not particularly high and there is only one Class~III YSO. In the center of the ridge, the dark cloud 338.8+16.5 is associated with recent star formation \\citep{96tachi,yancy}; in this area a coherent velocity gradient of about 1\\,\\kmspc\\ runs across the ridge. The SE end of the ridge is complex, with YSOs, enhanced line width and integrated intensity on the NW side, and column density (and possibly temperature) peaking to the SE. The enhanced line widths and velocity gradient in B\\,228 are consistent with a dynamical interaction between Lupus~I and the \\ion{H}{1} shell around the Upper-Sco subgroup of Sco-Cen \\citep{92degeus,01tachi}.  To the north of Lupus~III, the small cloud Lupus~III\\,N has similar characteristics to the bulk of the other clouds, albeit with broader lines and significant velocity structure. Lupus~III itself contains a compact CO peak which is probably heated by the nearby HAe/Be stars HR\\,5999 and 6000. The gas in this clump contains sufficiently dense gas to thermalize the CO 4--3 transition ($n_c\\sim 3\\times 10^4$\\,\\percc), but not to thermalize the \\hthcop\\ transition mapped by \\citet{07tachi}, which would require $n\\ga 10^5$\\,\\percc. The rest of Lupus~III seems to have quite similar physical conditions to those in the rest of the clouds, and shows no particular sign of being affected by the nearby OB subgroups.  Lupus~IV contains peaks of very high column density \\citep{mips} associated with slightly warmer gas temperatures (10--12\\,K). These temperatures are estimated from the optically-thick CO 4--3 transition, which is strongest around the extinction cores, suggesting significant external heating. The average line width of \\thco\\ 2--1, however, is significantly lower than those of Lupus~I and III. Lupus~IV faces the Upper-Cen-Lup subgroup of Sco-Cen and \\citet{02my} suggested that it is influenced by the OB stars; the Upper-Cen-Lup \\ion{H}{1} shell passed by the Lupus clouds a few Myr ago, so Lupus~IV is more likely to be affected by the radiation field from the OB stars.  Despite the basic similarities in their physical conditions, the three clouds have significant differences: Lupus~I appears to be strongly affected by external thermal and dynamical influences from the nearby Upper-Sco OB association, and does not display widespread star formation. Lupus~III shows no sign of external influence --- parts of the cloud are heated internally by its own young stars. Lupus~III\\,N seems entirely quiescent, yet has a large average line width. Lupus~IV has the greatest column density and the narrowest average line width, has almost no star formation as yet, and may be heated externally by the Upper-Cen-Lup OB association.  A detailed spatial comparison of CO and extinction maps will yield more accurate estimates of the physical conditions of the Lupus clouds, as well as mapping the variation in the \\thco-to-\\av\\ ratio. Mapping of additional CO transitions is crucial to the understanding of the clouds: more optically-thin lines (\\ceo\\ and even $\\mathrm{C^{17}O}$) are particularly important. Maps of the Upper-Sco \\ion{H}{1} shell with comparable resolution to the Lupus maps are required to look for more definitive signs of interaction between the shell and the molecular clouds. More sophisticated radiative transfer calculations are beyond the scope of this work, but are required to use the CO 4--3 emission as a proper constraint on the physical conditions of the gas, and hence to understand the structure of the Lupus clouds and how they are affected by the local environment."
    },
    "0910/0910.1505_arXiv.txt": {
        "abstract": "We develop a method to infer or rule out the presence of an atmosphere on a tidally-locked hot super Earth. The question of atmosphere retention is a fundamental one, especially for planets orbiting M stars due to the star's long-duration active phase and corresponding potential for stellar-induced planetary atmospheric escape and erosion. Tidally-locked planets with no atmosphere are expected to show a Lambertian-like thermal phase curve, causing the combined light of the planet-star system to vary with planet orbital phase.  We report {\\it Spitzer} 8$\\mu$m IRAC observations of GJ\\,876 taken over 32 continuous hours and reaching a relative photometric precision of $3.9 \\times 10^{-4}$ per point for 25.6~s time sampling. This translates to a 3-$\\sigma$ limit of $5.13 \\times 10^{-5}$ on a planet thermal phase curve amplitude. Despite the almost photon-noise-limited data, we are unable to conclusively infer the presence of an atmosphere or rule one out on the non-transiting short-period super Earth GJ\\,876d. The limiting factor in our observations was the miniscule, monotonic photometric variation of the slightly active host M star, because the partial sine-wave due to the planet has a component in common with the stellar linear trend.  The proposed method is nevertheless very promising for transiting hot super Earths with the {\\it James Webb Space Telescope} and is critical for establishing observational constraints for atmospheric escape. ",
        "introduction": "Super Earths are a recently discovered class of exoplanets, loosely defined to be 10~$M_{\\oplus}$ or less. Since these planets have low masses, they likely consist substantially of rocky material, making them the first analogs of terrestrial planets in our solar system. The search for super Earths around G through M stars continues \\citep[e.g.,][]{bagl2003, boru2003, butl2004, bonf2007, endl2008, nutz2008, mayo2008, mayo2009}, with many discoveries of transiting super Earths \\citep{roua2009} anticipated in the next few years.  A fundamental question about rocky planets is whether or not they have an atmosphere. The atmosphere can regulate the temperature and protect the surface from harmful solar or cosmic rays, issues central to surface habitability.  Extreme ultraviolet (EUV) radiation, winds, and coronal mass ejections \\citep[CMEs; ][]{lamm2007} from the host star are the main mechanisms for driving atmospheric escape. M stars have a much longer active phase compared to solar-like stars, during which a star's EUV, winds, and CMEs are strong \\citep{scal2007, west2008}.  The issue of atmospheric retention, therefore, is especially relevant for super Earths in short-period orbits around M stars \\citep[e.g.,][]{lamm2007}.  We are motivated to observationally determine whether or not a tidally-locked close-in super Earth orbiting an M star has an atmosphere.  The goal is to observe the planet and star in combined light and search for thermal phase variations peaking at the substellar point, indicative of a bare rock planet with no atmosphere.  The idea to study transiting hot super Earths in this way was briefly mentioned in \\citet{sels2004} and \\citet{nutz2008}. On a related topic, \\citet{gaid2004} and \\citet{sels2004} developed the same idea for non-transiting terrestrial planets orbiting at the Earth's semi-major axis from a sun-like star. The idea would be to use a Terrestrial Planet Finder/Darwin-type interferometer \\citep[e.g.,][and references therein]{cock2009, laws2008} to null out the star light to directly image the planet over the course of its orbit.  \\citet{gaid2004} explored light curves of planets with and without atmospheres and oceans, for a variety of obliquities, eccentricies, and viewing geometries. Combined light phase curves of about half a dozen hot Jupiter transiting planets have been obtained with the {\\it Spitzer Space Telescope} \\citep[e.g.,][]{harr2006, cowa2007, knut2007, knut2009}. Even though our first proposal to observe a hot super Earth phase curve was in 2005, three attempts at observations were needed to reach a reasonable S/N and to mitigate instrument systematics.  We begin in \\S2 with a description of the motivation and background to detect an atmosphere on a tidally-locked super-Earth. In \\S3 we describe {\\it Spitzer} IRAC observations of the planet-hosting M dwarf star GJ~876A and report on upper limits to inferring an atmosphere on the short-period super Earth GJ~876d. In \\S4 we discuss prospects for the method to infer or rule out an atmosphere on tidally-locked super-Earths, including transiting exoplanets. ",
        "conclusions": "We reported nearly continuous 32 hours of Spitzer IRAC data of GJ876d. The goal was to infer whether or not GJ~876d has an atmosphere, because a tidally-locked planet with no atmosphere is expected to show a Lambertian-like thermal phase curve.  We are unable to either conclusively infer the presence of an atmosphere or rule one out. We can present inclination- and albedo-dependent statements of what kind of planet atmospheres are ruled out if GJ876d is a planet of relatively low average density for a solid planet (see Figure 2).  The limiting factor in our data analysis is variation of the stellar flux. The presence of a slow linear drift increases the size of an atmosphereless planet that can hide in the data. Even though GJ 876d is a quiet M star, it still shows variability on a rotational time scale of 96.7 days \\citep{rive2005}.  The serendipitous detection of a flare (Figure~1) is further evidence for activity on a very slightly active star.  We emphasize that the flare itself should not impede efforts to extract the planet signal, because the flare has a specific shape and a relatively short duration.  The detection of infrared microflares on M stars may provide a new way to study M star activity not available from the ground.  Ground-based visible photometric monitoring of the star during the IR space-based measurements, combined with the rotation rate of the star and a star spot model (area and temperature of spots) would help to subtract any stellar variability in the light curve. Additionally, space-based IR monitoring before and after the planet observations would extend the data baseline to help precisely define the stellar variability. Future observations should also use continuous observations of the planet-star system, not just for maximizing S/N, but also to avoid the detrimental ``ramp'' effects that would be caused by observations of other targets of random brightnesses. Monitoring of a planet over two or more orbital periods would also help disentangle a linear stellar variability component.  Higher S/N would also mitigate effects of stellar variability. Our upper limit was degraded because of overlap between a component of the planet's variation and the stellar variability.  We fit to the data with a ``degraded template''.  Having more photons would help to get more precision, and tighter limits than using such a degraded template. We caution that with more photons we might encounter another aspect of stellar variability at a lower level. Our observations may be the first to indicate the difficulty even low levels of stellar variability have on observational studies in the combined-light of the planet and star, for any kind of exoplanets orbiting an M star.  Knowing the inclination of GJ\\, 876d would help in interpreting our data. \\citet{rive2005} provide an inclination for the GJ\\,876 3-planet system of $\\sim$~50 degrees, based on stability arguments for the two $\\sim$Jupiter-mass planets GJ~876b and GJ~876c, assuming that the lower mass GJ 876d is also coplanar.  {\\it Hubble Space Telescope} Fine Guidance Sensor observations \\citep{bene2002}, in contrast to the theoretical simulations, give an inclination of GJ 876b of 84$\\pm6$ degrees. Even with improved observations and theoretical simulations, the inclination of GJ 876d may remain unknown, unless we adopt the dynamical assumption that it should be co-planar with its more massive planet siblings.  Transiting planets provide a much better opportunity for observations to infer whether or not an atmosphere exists from a thermal phase curve compared to a planet with unknown orbital inclination. Because a transiting planet's orbital inclination is near 90 degrees, the planet-star flux contrast is maximized for a bare rock planet. The radius of a transiting planet is known, enabling a specification of the albedo without resort to mass-radius interior models. (While non-transiting planets with known inclination would still help with interpreting the results, the unknown radius remains a complicating factor.).  We can use Figure~2 to understand what kind of limits one could put for comparable data for a transiting planet. If GJ 876d were a transiting planet at 90 degrees inclination with no evident thermal phase variation, we would infer an atmosphere on GJ~876d for almost any interior composition. The exception is planets denser than a planet with a Mercury-like interior composition (approximately 70\\% iron core and 30\\% silicate mantle); slightly higher S/N data would be needed. For a planet with a Mercury-like composition, the albedo upper limit is 0.15. Most bare rocky bodies in our solar system have albedos lower than this value Mercury's geometric albedo is 0.1 and the Moon's is 0.12). One exception is Io, Jupiter's atmosphereless moon, with a high geometric albedo \\citep[0.63][]{cox2000} due to fresh solid sulfur deposits from episodic volcanic outbursts, themselves caused by tidal friction with Jupiter. Although many of the icy satellites of the outer planets also have very high geometric albedos \\citep[0.63][]{cox2000}, an ice-covered surface is not expected for hot planets like GJ 876d and others at very close planet-star separations. Data with S/N $\\sim$3 times higher than our data set reported here would rule out such high albedos even for the smallest plausible planets.  We have discussed how to infer the presence of an atmosphere on an exoplanet. Can an atmosphere be ruled out based on detection of a thermal phase curve? In principle a planet with a thick atmosphere having inefficient heat redistribution \\citep[e.g.,][]{harr2006} could produce a signal similar to a bare rock. Measuring the phase of the orbit variation to high precision is needed; a bare rock must peak very close to phase 0.5, but a planet with a thick atmosphere should have a detectable phase shift of the thermal maximum, even for relatively inefficient heat transport. Discriminating between a planet with a very thin atmosphere and a bare rock may be more difficult, and requires further study.  Observation of the combined light of the planet and star as a function of phase may also reveal the presence of a relatively strong planetary magnetic field. For short-period planets, the planet and star magnetic fields could interact to produce a phase curve peaking at phases near 180 degrees with a period equal to that of the planet's orbit, as has been detected for a few hot Jupiters \\citep{shko2005, dona2008, shko2008}. The question of super Earth magnetic fields is a serious one. Without a protective magnetic field the atmosphere might be further eroded by stellar wind and high energy particles. A detection of a relatively strong magnetic field would indicate the presence of a partially liquid interior, aiding interpretation of interior composition (and even further interesting for short-period planets under strong tidal interaction.)  We note that Mercury is not tidally-locked to the sun, (i.e., in a 1:1 resonance). Instead Mercury, is in a 3:2 resonance, rotating three times for every two orbits about the sun. Models show a low probability for a planet on a circular orbit to become trapped in the 3:2 spin-orbit resonance compared to the 1:1 resonance, unless it goes through a chaotic evolution of an eccentric orbit as Mercury likely has \\citep{corr2004}. Planets on circular orbits are therefore more certainly tidally-locked than eccentric planets which may have had the opportunity to be captured into higher order spin-orbit resonances.  The {\\it James Webb Space Telescope}, with its 58 times more collecting area than {\\it Spitzer} will be able to make high S/N observations of short-period super Earths. M stars are good targets, because they are bright at IR wavelengths and the planet-star radius ratio is high.  In conclusion, inferring an atmosphere on a short-period, tidally-locked super Earth that lacks a thermal phase curve is possible for a transiting planet of any interior composition and albedo, provided that the data set has higher S/N by about a factor of 3 than ours. During the {\\it JWST} era we anticipate the birth of studies of atmosphereless bodies and the concomitant understanding of atmospheric escape."
    },
    "0910/0910.5110_arXiv.txt": {
        "abstract": "The lifetime of protoplanetary discs is intimately linked to the mechanism responsible for their dispersal. Since the formation of planets within a disc must operate within the time frame of disc dispersal it is crucial to establish what is the dominant process that disperses the gaseous component of discs around young stars. Planet formation itself as well as photoevaporation by energetic radiation from the central young stellar object have been proposed as plausible dispersal mechanisms. There is however still no consensus as what the dominant process may be. In this paper we use the different metallicity dependance of X-ray photoevaporation and planet formation to discriminate between these two processes. We study the effects of metallicity, $Z$, on the dispersal timescale, $t_{\\rm phot}$, in the context of a photoevaporation model, by means of detailed thermal calculations of a disc in hydrostatic equilibrium irradiated by EUV and X-ray radiation from the central source. Our models show $t_{\\rm phot} \\propto Z^{0.52}$ for a pure photoevaporation model. By means of analytical estimates we derive instead a much stronger {\\it negative} power dependance on metallicity of the disc lifetime for a dispersal model based on planet formation.  A census of disc fractions in lower metallicity regions should therefore be able to distinguish between the two models. A recent study by Yasui et al. in low metallicity clusters of the extreme outer Galaxy ([O/H]$\\sim$-0.7dex and dust to gas ratio of $\\sim$0.001) provides preliminary observational evidence for shorter disc lifetimes at lower metallicities, in agreement with the predictions of a pure photoevaporation model. While we do not exclude that planet formation may indeed be the cause of some of the observed discs with inner holes, these observational findings and the models and analysis presented in this work are consistent with X-ray photoevaporation as the dominant disc dispersal mechanism.  We finally develop an analytical framework to study the effects of metallicity dependent photoevaporation on the formation of gas giants in the core accretion scenario. We show that accounting for this effect strengthens the conclusion that planet formation is favoured at higher metallicity. We find however that the metallicity dependance of photoevaporation only plays a secondary role in this scenario, with the strongest effect being the positive correlation between the rate of core formation and the density of solids in the disc. ",
        "introduction": "It is not yet established what is the dominant process that disperses the gaseous component of the discs around young stars. What is clear observationally is that evidence of gas accretion onto stars, together with infrared/submm  diagnostics associated with small dust grains entrained in the gas, disappear in young stars at an average age of a few Myr (Haisch, Lada \\& Lada 2001). This figure needs to be interpreted as an {\\it average} figure since clearly individual stars may lose their discs on timescales that differ from this by at least a factor three (Armitage, Clarke \\& Palla 2003). It is equally clear, based on the relatively few systems observed in a state of transition between disc possessing and discless status, that the timescale for disc dispersal (specifically the timescale over which regions of the disc are optically thin in the infrared) is much shorter than the overall disc lifetime (Skrutskie et al. 1990, Kenyon \\& Hartmann 1995, Duvert et al. 2000). The fact that transition discs are relatively rare (constituting around $10 \\%$ of the population of young stars with discs) is however an observational hindrance when it comes to establishing the mechanism for disc clearing: transition discs appear to be a rather heterogeneous class of objects and it is hard to classify their diversity when the number of well studied objects in nearby star forming regions is still in single figures. An important feature of many transition discs, however, is that their spectral energy distributions are best fit by inner holes in the disc (these holes are not necessarily devoid of dust, but there is a large contrast in surface density between the outer disc and inner cavity).\\footnote{This situation may not extend to discs around lower mass stars, as there are recent claims that discs in M stars may instead pass through an extended phase in which their dust content is homogeneously depleted at all radii (Currie et al. 2009).}  Although an obvious mechanism for the disappearance of  dust diagnostics in discs is the simple coagulation of grains into entities that are large compared with the wavelength of observation (Dullemond \\& Dominik 2005), this scenario does not of itself explain why, in general, there is a correlation between the disappearance of dust and gas accretion diagnostics. Moreover, it is not obvious why dust coagulation should create a well defined inner hole structure rather than homogeneous depletion at all radii (although models combining grain growth with photophoresis - once the inner disc is optically thin to visible light in the radial direction  - are promising in this regard; Krauss et al. 2007).  The two leading disc dispersal mechanisms that satisfy these considerations are the formation of giant planets and photoevaporation. Both these mechanisms create an inner hole in the disc and also disrupt the accretion flow onto the star to some extent\\footnote{ More or less completely in the case of photoevaporation (Clarke et al 2001, Alexander, Clarke \\& Pringle 2006b) and to a variable extent in the case of planet formation, depending on the mass of the planet (Lubow et al. 1999, Rice et al. 2003).}. In the case of photoevaporation, inner hole creation is succeeded by the rapid clearing of the outer disc thereafter. In the case of planet formation, clearing of the outer disc requires further planet formation at larger radii, i.e. the relative scarcity of transition discs argues that planet formation at one radius is rather rapidly followed by a wave of successive planet formation events in the outer disc ( Armitage \\& Hansen  1999). Some attempts have been made to classify whether individual transition discs are likely to be  generated by planet formation or photoevaporation (Najita et al. 2007, Alexander \\& Armitage 2007,  Alexander 2008, Cieza et al. 2008, Kim et al. 2009) based on estimates of the accretion rate and disc mass. These analyses are however based on extreme ultraviolet (EUV) photoevaporation models (Hollenbach et al. 1994, Clarke et al. 2001, Alexander, Clarke \\& Pringle 2006a,b) which predict low accretion rates and disc masses in transition objects, in contrast to many of the observed systems. More recent models based on X-ray photoevaporation (Ercolano et al 2008, 2009, Owen et al. 2009) predict higher accretion rates and disc masses during transition and thus somewhat muddy the observational distinction between the two classes of disc dispersal mechanism. Nevertheless it would appear clear that no single mechanism can explain all observed inner hole systems: for example, photoevaporation is incompatible with the predominantly high accretion rates in the sample of Kim et al. (2009) whereas a planet is perhaps an unlikely candidate for many of the objects contained in Cieza et al. (2008), which have very large holes and low accretion rates.   \\begin{figure*} \\begin{center} \\begin{minipage}{20cm} \\includegraphics[width=9.0cm]{zmdot.eps} \\includegraphics[width=9.0cm]{ztau.eps} \\end{minipage} \\caption[]{The metallicity dependance of mass loss rates (left panel) and dispersal timescales (right panel) of protoplanetary discs photoevaporated by X-ray and EUV from the central star. {\\it Right} Model data points (asterisks) obtained from our photoionisation modeling; the dashed line is a least square fit to these data. {\\it Left:} The asterisks show the photoevapoaration timescales obtained substituting the points on the right panel into equation A10 and normalising to 2~Myr for solar metallicities. The dashed line shows the relation given in equation 4 for p~=~1.} \\label{f:mods} \\end{center} \\end{figure*}   At this point it is perhaps worth noting - despite the obvious interest of observed transition objects - that they do not {\\it necessarily} hold the key to the process that disperses the {\\it majority} of discs around young stars. This is because objects observed in transition can in principle be a mixture of stars that spend a significant time in this state (and must therefore be a minority of all young stars on statistical grounds) and a general population of `typical' young stars that must pass through such a transition quickly. (The argument that some fairly abrupt end to a disc's lifetime is required is well illustrated by considering what would happen if a disc merely continued to accrete onto the central star as a result of viscous evolution - the observed accretion rates, ages and disc masses of T Tauri stars would imply that integrated forward in time they would remain optically thick for $> 100$Myr and then would linger for a comparable period as partially cleared systems). Therefore  - regardless of what process explains the bulk of observed transition discs - we also need to decide what is the mechanism that terminates the lifetime of the majority of disc bearing young stars.  In this paper we therefore take a different approach to discriminating between planet formation and photoevaporation as the prime disc dispersal mechanism, by instead considering the {\\it metallicity dependence} of these two processes. Although most studies of circumstellar discs have currently been undertaken in regions with a rather low range of metallicities, the advent of the Herschel Space Observatory and later the James Webb Space Telescope raises the prospect of being able to conduct disc censuses in more distant regions whose metallicity differs considerably from that of star forming regions in the solar neighbourhood. Such studies are already possible in the Extreme Outer Galaxy (EOG) with current instrumentation as demonstrated by the recent work of Yasui et al. (2009) which we will discuss in more detail later. Our focus here will not be on the nature of the (in any case rare) transition objects, but will instead centre on how the {\\it mean disc lifetime is expected to vary with metallicity in these two scenarios}. To this end we combine simple models for disc viscous evolution  with (metallicity dependent) photoevaporation and also with semi-analytic prescriptions for planet formation (Section 2). We will show that in the photoevaporation model the disc lifetime is a mildly increasing function of metallicity, resulting from the rather higher photoevaporation rates in the case of low metallicity gas for which opacities are lower and line cooling less efficient. On the other hand, the disc lifetime is a strongly decreasing function of metallicity in the case that disc clearing is initiated by planet formation. This is simply because at  higher metallicity the assumed higher surface density of planetesimals in the disc encourages the more rapid formation of a gas giant planet, specifically because it accelerates the formation of a rocky core of the critical mass required to initiate the accretion of a gaseous envelope. This is borne out, for example, by the hydrodynamical models of gas giant formation of Pollack et al. (1996) and Hubickyj et al. (2005)\\footnote{Note that this conclusion is based on an assumed linear relationship between the density of planetesimals and the metallicity; this conclusion would be only strengthened if one takes into account the recent suggestion of  Anders et al 2009  that the efficiency of planetesimal formation by the streaming instability should increase steeply at higher metallicity.}. This positive link between metallicity and planet formation has been noted by a number of authors (e.g. Ida \\& Lin 2004b) as a possible explanation of the observed increase in the frequency of giant planets as a function of metallicity (Santos, Israelian \\& Mayor 2000, 2001, 2004; Gonzalez et al. 2001; Sadakane et al. 2002; Heiter \\& Luck 2003; Laws et al. 2003; Fischer \\& Valenti 2005; Santos et al. 2002 and references therein) but has not been previously linked to the issue of disc lifetimes. In Section 3, we revisit the connection between the frequency of giant planets and metallicity in the case of a hybrid model where planet formation takes place in the context of a nebula that is subject to metallicity dependent photoevaporation of the disc gas. Since photoevaporation is more effective at low metallicities and thus reduces the disc lifetime, this strengthens the conclusion that planet formation is favoured at high metallicity. Nevertheless, the  additional ingredient of metallicity dependent photoevaporation is found to be a second order effect in determining the positive correlation between planet frequency and metallicity. We present our conclusions in Section 4. ",
        "conclusions": "We have shown that the study of disc lifetimes in regions of different metallicities can be used as a powerful discriminant between the two currently leading models of disc dispersal -- photoevaporation and planet formation. By means of detailed thermal and photoionisation calculations we have determined that a disc's lifetime against photoevaporation is a shallow positive power of metallicity, $t_{\\rm phot} \\propto Z^{0.52}$. This metallicity dependance is specific to a photoevaporation mechanism driven mainly by X-ray radiation as it relies on the significant reduction of the gas opacities with metallicity. We have also shown that, on the contrary, a disc dispersal mechanism based on planet formation yields disc lifetimes that are a strong negative power of $Z$. Therefore, a census of disc fractions in regions of lower metallicities compared to the solar neighbourhood will be crucial to determine which is the dominant mechanism responsible for the rapid demise of protoplanetary discs.  Recent observations of embedded clusters in Cloud 2 of the Extreme Outer Galaxy (EOG) presented by Yasui et al. (2009, YS09) indeed find shorter disc lifetimes for this metal poor environment ([O/H]$\\sim$-0.7 dex), compared to the solar neighbourhood. These results are in agreement with the predictions of an X-ray+EUV photoevaporation model, and argue against planet formation as the dominant dispersal mechanism.  YS09 quote a mass detection limit of 0.1~M$_{\\odot}$ for Cloud2 in the EOG using the 8.2m {\\it Subaru} telescope, similar to values obtained for embedded clusters in the solar neighborhood with smaller (2-4m class) telescopes, and thus argue that a comparison of the disc fractions in embedded clusters in the EOG with those of clusters in the solar neighbourhood is justified. Further observations aimed at determining disc fractions in low metallicity regions of various ages would be very useful to confirm these important results.  We finally show that, in the context of giant planet formation in the core accretion scenario, the effect of metallicity dependent photoevaporation is to strengthen the conclusion that planet formation is favoured in high metallicity environments since the lifetime of the disc against photoevaporation ($t_{\\rm phot}$) is a positive function of Z. This effect, however, only plays a secondary role: the main process that favours planet formation at high metallicity is simply the faster core growth in the case of a high surface density of solids in the disc."
    },
    "0910/0910.2730_arXiv.txt": {
        "abstract": " ",
        "introduction": "The matter-free spacetime is strictly flat, more precisely Ricci-flat, unless its fabric is endowed with an intrinsic curvature $\\Lambda_0$. The ensuing spacetime curvature is governed by \\begin{eqnarray} \\label{field0} R_{\\alpha \\beta}\\left(\\Gamma\\right) - \\frac{1}{2} g_{\\alpha \\beta} R\\left(\\Gamma\\right) = - \\Lambda_0\\, g_{\\alpha \\beta} \\end{eqnarray} where $R_{\\alpha \\beta}$ is the Ricci tensor, $R$ the Ricci scalar, and \\begin{eqnarray} \\label{levi-civita} \\Gamma^{\\lambda}_{\\alpha \\beta} = \\frac{1}{2} g^{\\lambda \\rho} \\left( \\partial_{\\alpha} g_{\\beta \\rho} + \\partial_{\\beta} g_{\\rho \\alpha} - \\partial_{\\rho} g_{\\alpha \\beta}\\right) \\end{eqnarray} is the Levi-Civita connection. The constant curvature term $\\Lambda_0$, Einstein's cosmological constant \\cite{einstein1}, represents an incalculable source of curvature, independent of the entirety of matter and forces but gravity. As it stands, it is a constant of Nature the value of which is to be determined empirically.  In the presence of matter and radiation, the matter-free gravitational field equations (\\ref{field0}) change to \\begin{eqnarray} \\label{field1} R_{\\alpha \\beta}\\left(\\Gamma\\right) - \\frac{1}{2} g_{\\alpha \\beta} R\\left(\\Gamma\\right) = - \\Lambda_0\\, g_{\\alpha \\beta} + 8 \\pi G_N T_{\\alpha \\beta} \\end{eqnarray} which is nothing but augmentation of (\\ref{field0}) by the energy-momentum tensor $T_{\\alpha \\beta}$ of matter and radiation \\cite{einstein2}. The prime idea in passing from (\\ref{field0}) to (\\ref{field1}) is that gravitational field is described by Poisson equation in the Newtonian limit.  In general, $T_{\\alpha \\beta}$ is computed from the quantum effective action for a given background metric $g_{\\alpha \\beta}$ which necessarily metamorphoses into a dynamical variable through (\\ref{field1}). On general grounds, energy-momentum tensor assumes the generic form \\begin{eqnarray} \\label{en-mom} T_{\\alpha \\beta} = - \\texttt{E}\\, g_{\\alpha \\beta} + {\\texttt{t}}_{\\alpha \\beta} \\end{eqnarray} where $\\texttt{E}$ stands for the energy density of the vacuum state, and  ${\\texttt{t}}_{\\alpha \\beta}$ does for the contributions of matter and radiation, collectively. Replacement of this decomposition of $T_{\\alpha \\beta}$ into (\\ref{field1}) gives rise to an effective cosmological constant \\begin{eqnarray} \\label{lameff} \\Lambda_{\\texttt{eff}} = \\Lambda_0 + 8 \\pi G_N \\texttt{E} \\end{eqnarray} which must nearly saturate the expansion rate of the Universe \\begin{eqnarray} \\label{cond} \\Lambda_{\\texttt{eff}} \\simlt H_0^2 \\end{eqnarray} since the non-vacuum mass contained in $\\texttt{t}_{\\alpha \\beta}$ is a small fraction of the critical density. A number of independent observations \\cite{astro1,astro2,astro3,astro4,astro5,astro6,astro7,astro8} have determined $H_0$ to measure approximately $73.2\\ {\\rm Mpc}^{-1}\\, {\\rm s}^{-1}\\, {\\rm km}$. This observational result provides an experimental determination of $\\Lambda_{\\texttt{eff}}$: \\begin{eqnarray} \\label{experi} \\Lambda_{\\texttt{eff}}^{\\texttt{exp}} \\simeq 8 \\pi G_N \\texttt{E}^{\\texttt{exp}} \\end{eqnarray} where $\\texttt{E}^{\\texttt{exp}} \\simeq 3.25\\times 10^{-47}\\ {\\rm GeV}^4$  \\cite{astro1,astro6,astro7} which, by just numerical coincidence, equals approximately $m_{\\nu}^4$ with $m_{\\nu} \\simeq 10^{-3}\\ {\\rm eV}$ being the neutrino mass.  If it were $\\Lambda_0$ not $\\Lambda_{\\texttt{eff}}$, the bound (\\ref{cond}) would furnish, through the observational value of $H_0$ quoted above, an empirical determination of $\\Lambda_0$, as for any other fundamental constant of Nature. The same does not apply to $\\Lambda_{\\texttt{eff}}$, however. The reason is that the vacuum energy density $\\texttt{E}$, equaling the zero-point energies of quantum fields plus enthalpy released by various phase transitions, turns out to be characteristically much larger than $\\Lambda_{\\texttt{eff}}^{\\texttt{exp}}/ 8 \\pi G_N$. That this is the case can be illustrated by considering, for instance, the electron which weighs next to neutrinos. The electron loop gives sizeable contributions to $\\texttt{E}$. The smallest energy density it gives is electron mass per its cubic Compton wavelength, and it is already much larger than $\\Lambda_{\\texttt{eff}}^{\\texttt{exp}}/ 8 \\pi G_N$. Much grosser than this is that the electron loop contributes to $\\texttt{E}$ by additional terms growing quadratically and quartically with the ultraviolet cutoff. Consequently, the known, experimentally confirmed fields and forces down to the terascale, $M_{W} \\sim {\\rm TeV}$, are expected to induce a vacuum energy density of order $M_{W}^4$  -- equaling the sum total of energies deposited by parton-hadron and electroweak phase transitions and by the quantal zero-point energies of fields. This energy density, by itself, gives rise to a centimeter-size Universe unless it is neutralized by the $\\Lambda_0$ contribution in (\\ref{lameff}) -- a severe tuning of numbers up to at least sixty decimal places. This immense tuning worsens if the standard model of strong and electroweak interactions extends beyond Fermi energies without a suppression mechanism for $\\texttt{E}$. One thus concludes that, enforcement of $\\Lambda_{\\texttt{eff}}$ to obey (\\ref{cond}) gives rise to the biggest naturalness problem plaguing both particle physics and cosmology -- the cosmological constant problem. The CCP is a highly inextricable perplexity a resolution of which is likely to be found outside the framework set by the gravitational field equations (\\ref{field1}).  Over the decades, since its first solidification in \\cite{ccp1,ccp2}, the CCP has been approached by putting forth various proposals and interpretations, as listed and critically discussed in \\cite{nobbenhuis1,nobbenhuis2} (see also the review volumes \\cite{review1,review2,review3,review4,review5} and references therein). They each involve necessarily a certain degree of speculative aspect in regard to going beyond (\\ref{field1}). These aspects involve postulating novel symmetry arguments, relaxation mechanisms, modified gravitational dynamics and statistical interpretations (see \\cite{1,2,3,4,5,6,7,8,9,10,11,12,13,14,15,16,17,18,19,20,21x,22,23,24,25,26,27,28,29,30,31,32,33,34,35,36,37,38,39,40,41,42,43,44,45,46,47,48,48p,49} for a partial list of recent work) as discussed in \\cite{nobbenhuis1,nobbenhuis2}. Excepting the nonlocal, acausal modification of gravity implemented in \\cite{nima1,nima2} (see also \\cite{33}) and the anthropic approach \\cite{anthro}, most of the solutions proposed for the CCP seem to overlook the already existing vacuum energy density ${\\mathcal{O}}\\left({\\rm TeV}^{4}\\right)$ induced by the known physics down to the terascale \\cite{weinberg1,weinberg2,weinberg3}. Indeed, any resolution of the CCP, irrespective of how speculative it might be, must, in the first place, provide an understanding of how this existing energy component to be tamed.  Having stated the problem, we are at the stage to lay out the germ of the mechanism to be proposed in the present work. The alleged mechanism rests on finding a sensible answer to the question: {\\it Can one excogitate a way, different than in (\\ref{field1}), of incorporating $T_{\\mu \\nu}$ into the matter-free gravitational field equations (\\ref{field0}) while keeping all the successes of the Einstein field equations for non-vacuum sources $\\texttt{t}_{\\alpha \\beta}$ yet naturalizing the effects of the vacuum energy $\\texttt{E}$ ?} The answer is affirmative. The method, as will be described in Sec. 2 below, involves incorporation of matter and radiation into (\\ref{field0}) by replacing the metric $g_{\\alpha \\beta}$ by a general tensor field $\\mathbb{T}_{\\alpha \\beta}$ -- to be related to the energy-momentum tensor $T_{\\alpha \\beta}$. In Sec. 3 it will be shown that the dynamics in Sec. 2 follow from an action principle. Sec. 4 is devoted to a critical discussion of the method  analyzed in Sec. 2. In Sec. 5 we conclude. ",
        "conclusions": "Since its first solidification in \\cite{ccp1}, the CCP has caused a huge literature avalanche. The problem is a deeply perplexing one, and a resolution seems unlikely to be found in the framework of quantum field theory and General Relativity. The literature is widespread in terms of topics, scopes and methods of the attempts at understanding the CCP. To mention a few of the main strategies, one notes various works utilizing symmetry-based arguments \\cite{nobbenhuis1,nobbenhuis2,review4,5,6,8,11,12,13,15,16,26,34,38,42,49}, relaxation mechanisms \\cite{nobbenhuis1,review1,review2,review3,review4,review5,7,9,10,17,22,23,28,29,30,33,46,47,48,48p}, and modified gravitational theories \\cite{nobbenhuis1,review1,review5,1,2,4,14,18,19,25,31,32,36,37,39,40,41,43,44,45,nima1,nima2}. Given the well-understood part of Nature down to the terascale, a fundamental measure of the validity of any proposition is its capability to degravitate the vacuum energy density (${\\cal{O}}\\left({M}_{W}^4\\right)$ or higher) induced by the matter sector. Therefore, each attempt at solving the CCP can be judged upon this minimal requirement plus its validity and generality. Quite expectedly, any change designed for solving the CCP must give rise to novel effects in quantum field theoretic and gravitational contexts, and it is via these effects that one can hint in the true mechanism behind the tiny cosmological constant measured \\cite{astro1,astro2,astro3,astro4,astro5,astro6,astro7,astro8}.   Compared to the existing literature, the present work brings up a novel, geometro-dynamical mechanism for a resolution  of the CCP. The mechanism, based on the six propositions stated and proven in Sec. 2, enables one to distinguish between Einstein's cosmological constant $\\Lambda_0$ and the vacuum energy deposited by the quantal matter and radiation. Thus, the vacuum energy of matter sector, instead of gravitating, facilitates the generation of the gravitational constant. If gravity is a classical phenomenon, as has been claimed for the de Sitter space generated by $\\Lambda_0$, then the proposed mechanism can be regarded to have naturalized the CCP since determination of $\\Lambda_0$, in isolation, from cosmological observations involves no fine-tuning at all. On the other hand, if gravity is quantized, then CCP gets revived due to graviton loops, and the mechanism proposed here is simply halted to offer any solution. In this particular case, all that the present method can do is to modify the nature of the CCP in that it becomes a naturalness problem pertaining to the gravitational sector, only. Quantum gravity becomes the main obstacle in searching for a way to suppress violent corrections from quantum fluctuations. However, if de Sitter gravity remains as a purely classical structure in accord with various claims then the mechanism proposed in this work gains a decisive status for the CCP.  The fundamental equations hypothesized are free of the CCP to start with. Reproduction of correct gravitational dynamics, in a way free of the CCP, has been accomplished by constructing a new geometry whose metric tensor is a linear, causal and nonlocal functional of the energy-momentum tensor of the matter and radiation. If the covariantly-constant part of this metric is large enough to facilitate a power series expansion, at the lowest order, gravitational field equations are obtained by an appropriate definition of the gravitational constant. Nevertheless, the original linear relation must be modified order by order in perturbation theory to reach Bianchi consistency. The subleading terms obtained this way contain nonlocal pieces whose effects can be crucial for description of Nature at super-Planckian energies where the aforementioned power series expansion necessarily breaks down. For physical phenomena at ordinary energies, however, the resulting dynamical equations are indistinguishable from the ones in General Relativity to an excellent approximation. In heart, the mechanism makes critical use of the dual nature of the metric tensor: It defines both the spacetime geometry and energy-momentum tensor of the vacuum."
    },
    "0910/0910.0445_arXiv.txt": {
        "abstract": "This contribution gives a brief overview of the theoretical ideas underlying our current understanding of the early Universe. Confronting the predictions of the early Universe models with cosmological observations, in particular of the cosmic microwave background fluctuations,  will improve our knowledge about the physics of the primordial Universe. ",
        "introduction": "Inflation is today the main theoretical framework to describe  the early Universe. In thirty years of existence,  inflation, in contrast with earlier competitors,   has survived the confrontation with  cosmological data, which have tremendously improved over the years.  Indeed, the fluctuations of the Cosmic Microwave Background (CMB) had not yet been measured when inflation was invented, whereas they give us today  a remarkable   picture of  the cosmological perturbations in the early Universe.  In the future, one can hope that even more precise observations will allow us to test  inflation further, and also to discriminate between the many different possible realizations of inflation.  This contribution discusses the basic ideas underlying inflation (many more details  can be found in e.g. \\cite{cargese}) and some more recent results. ",
        "conclusions": ""
    },
    "0910/0910.4167.txt": {
        "abstract": "{}{}{}{}{} % 5 {} token are mandatory  \\abstract % context heading (optional) % aims heading (mandatory) {} {This paper reports on H-band interferometric observations of \\object{Betelgeuse} made at the three-telescope interferometer IOTA. We image \\object{Betelgeuse} and its asymmetries to understand the spatial variation of the photosphere, including its diameter, limb darkening, effective temperature, surrounding brightness, and bright (or dark) star spots.} % (to look at signatures of convection?) % methods heading (mandatory) {We used different theoretical simulations of the photosphere and dusty environment to model the visibility data.  We made images with parametric modeling and two image reconstruction algorithms: MIRA and WISARD.} % results heading (mandatory) { We measure an average limb-darkened diameter of 44.28 $\\pm$ 0.15\\,mas with linear and quadratic models and a Rosseland diameter of 45.03 $\\pm$ 0.12\\,mas with a MARCS model. These measurements lead us to derive an updated effective temperature of 3600 $\\pm$ 66\\,K. We detect a fully-resolved environment to which the silicate dust shell is likely to contribute. By using two imaging reconstruction algorithms, we unveiled two bright spots on the surface of \\object{Betelgeuse}. One spot has a diameter of about 11 mas and accounts for about 8.5\\% of the total flux. The second one is unresolved (diameter $<$ 9 mas) with 4.5\\% of the total flux.} % conclusions heading (optional), leave it empty if necessary {Resolved images of \\object{Betelgeuse} in the H band are asymmetric at the level of a few percent. The MOLsphere is not detected in this wavelength range. The amount of measured limb-darkening is in good agreement with model predictions. The two spots imaged at the surface of the star are potential signatures of convective cells.} ",
        "introduction": "Located in the Orion constellation, \\object{Betelgeuse} ($\\alpha$ Orionis) is a red supergiant (hereafter RSG) of spectral type M2Iab. It is one of the brightest stars at optical wavelengths and has the second biggest angular diameter \\citep[$\\sim$43 mas,][]{2004A&A...418..675P} after R Doradus \\citep{1997MNRAS.286..957B}. Classified as semi-regular, it shows periodicity in its brightness changes, accompanied or sometimes interrupted by various irregularities \\citep{1984LNP...193..336G}. According to a recent reanalysis of a Hipparcos satellite dataset, its distance is now estimated at 197 $\\pm$ 45 pc \\citep{2008AJ....135.1430H}. The observations up to now have identified at least 7 components of the complex \\object{Betelgeuse} atmosphere: two outer shells, a dust environment, a chromosphere, a gaseous envelope, a molecular shell also known as MOLsphere \\citep{2000ApJ...538..801T}, and finally the photosphere. Some of these components are not symmetric, and some are overlapping.    \\subsection{Two outer shells} Two shells (S1 and S2) have been observed in absorption in $^{12}$C$^{16}$O and $^{13}$C$^{16}$O at 4.6 $\\mu$m by \\cite{1979ApJ...233L.135B} but their spatial extent was not directly determined. Phoenix 4.6 $\\mu$m spectra obtained by \\cite{2009AIPC.1094..868H} led to an estimation of the size for S1 and S2. An outer radius of $\\sim$4.5 and $\\sim$7 arcsec were derived, respectively,  although the latter is inconsistent with other measurements made by the CARMA interferometer.   \\subsection{A dusty environment} \\label{dustshell} A shell of dust was first detected with heterodyne interferometry at 11 $\\mu$m by \\cite{Sutton}. Quite a few estimates of the inner radius of \\object{Betelgeuse}'s dust shell have since been made. \\cite{Bester1991} used a model with an inner radius of 0.9\\,arcsec ($\\sim$45\\,R$_{\\star}$) to explain both their 11\\,$\\mu$m heterodyne interferometry (ISI) and older speckle observations by \\cite{Sutton} and \\cite{Howell1981}. In their spatially-resolved mid-infrared (mid-IR) slit spectroscopy, \\cite{Sloan1993} find no silicate emission within the central arcsecond around \\object{Betelgeuse}\\footnote{They also remark that this actually argues against a spherical distribution of the dust.}. \\cite{Danchi1994} find from 11.15\\,$\\mu$m ISI data that the inner radius must be $1.00\\pm0.05$\\,arcsec, i.e. roughly 50\\,R$_{\\star}$. But this result disagrees with later findings by \\cite{Skinner1997}, who claim an inner radius of not more than 0.5\\,arcsec. This environment of dust is also reported by \\cite{2007ApJ...670L..21T} as the origin of a visibility drop at short baselines of more than 40\\% in new ISI measurements.   \\subsection{Chromosphere} \\cite{1975MNRAS.172..277L} interpreted \\object{Betelgeuse} spectra and found that  the infrared excess could be explained by a chromospheric emission. More recently, the ultraviolet (UV) observations of the Hubble Space Telescope revealed a spatially resolved chromosphere measuring twice the size of the optical photosphere in the continuum \\citep{1996ApJ...463L..29G,1998AJ....116.2501U}.  \\subsection{The asymmetric gaseous envelope} The outer envelope of the O-rich \\object{Betelgeuse} \\citep[O/C=2.5,][]{1984ApJ...284..223L} has been the subject of extensive studies. VLA observations revealed an extended, cool ($\\sim$1000-3000 K) and irregularly shaped gaseous atmosphere of several stellar radii which coexists with the chromospheric gas \\citep{1998Natur.392..575L}. More recently \\cite{2009arXiv0907.1843K} found evidence for the presence of a bright plume, possibly containing CN, extending in the southwest direction up to a minimum distance of six photospheric radii. Such a plume could be formed either by an enhanced mass loss above a large upwards moving convective cell or by rotation, through the presence of a hot spot at the location of the polar cap of the star.   \\subsection{The MOLsphere} With the Stratoscope II experiment, strong water-vapor bands have been detected in the spectrum of \\object{Betelgeuse} \\citep{1965ApJ...141..116D} but at that time no explanation on their origin was given. More recently, Tsuji \\citep{2000ApJ...538..801T,2000ApJ...540L..99T} proposed the existence of a molecular layer at 0.45 $R_{\\star}$ above the photosphere, a MOLsphere, of effective temperature 1500 $\\pm$ 500 K  and composed of water vapor and CO. Using infrared interferometry, \\cite{2004A&A...418..675P} showed that K, L and 11.15 $\\mu$m data could be simultaneously modeled with a gaseous shell of temperature 2055 $\\pm$ 25 K located 0.33 $R_{\\star}$ above the photosphere while \\cite{2006ApJ...637.1040R} locate the water vapor inside the photosphere by analyzing high-resolution spectra. A model based on mid-infrared measurements showed that the MOLsphere could contain not only water vapor but also SiO and Al$_{2}$O$_{3}$ \\citep{Verhoelst2006,2007A&A...474..599P}. This latter component is a type of dust condensing at high temperature and on which SiO can be adsorbed. It is a nucleation site at high temperature for silicate dust.  This is a possible scenario to explain dust formation and to connect molecular layers and dusty environments.   \\subsection{The convective photosphere} \\cite{2004A&A...418..675P} modeled the photosphere of \\object{Betelgeuse} as a uniform disk of  3641 $\\pm$ 53 K temperature with a 42 mas diameter. For some wavelengths, the diameter can vary with time. At 11.15$\\mu$m, \\cite{2009ApJ...697L.127T} measure a systematic decrease of the  diameter of \\object{Betelgeuse} during the period 1993-2009 that might point toward long-term dynamics at its surface. Moreover, the more recent hydro-radiative simulations of RSGs \\citep{2002AN....323..213F} foresee the presence of convection cells whose size is comparable to the stellar radius, as suggested by \\cite{1975ApJ...195..137S}.  \\subsection{The spots observed on the surface of \\object{Betelgeuse}} %Because of its big angular size, \\object{Betelgeuse} surface has been observed with 10 meter telescopes. Thanks to the pupil masking technique, \\cite{1992MNRAS.257..369W, 1997MNRAS.291..819W}  and \\cite{1990MNRAS.245P...7B} have detected the presence of spots on the surface of \\object{Betelgeuse} that represent up to 15-20\\% of the observed flux in the visible. The most likely origin of these asymmetries is a signature of the convection phenomenon. However, the hypotheses of inhomogeneities in the molecular layers or the presence of a transiting companion are not ruled out \\citep{1990MNRAS.245P...7B}. These observed spots are few, 2 or 3, and their characteristrics change with time \\citep{1992MNRAS.257..369W}. Confirming previous results, \\cite{2000MNRAS.315..635Y} observed spots in the visible but show a fully centro-symmetric picture of \\object{Betelgeuse} in the near infrared (at 1.29 $\\mu$m). Lately \\cite{2007ApJ...670L..21T} also report the observation of an asymmetry at 11.15 $\\mu$m located on the southern edge of the disk of \\object{Betelgeuse}. In order to deepen our knowledge on RSGs, it is critical to observe spots at high angular resolution. Many unknowns remain such as their size, location, chemical composition, dynamical properties, lifetime and origin. \\vspace{3mm}     To investigate these questions, we report in this paper on interferometric observations obtained with the IONIC 3-telescope beam combiner at the IOTA interferometer in October 2005. Sect.~\\ref{obssect} describes the observations and the data reduction. In Sect.~\\ref{envsect}, the fully resolved environment is discussed. The physical diameter and limb darkening of \\object{Betelgeuse} are investigated in Sect.~\\ref{LDsect} . In Sect.~\\ref{imasect}, images of \\object{Betelgeuse} reconstructed with the MIRA and WISARD algorithms are compared to a parametric fit. Conclusions are gathered in Sect.~\\ref{conclusect}.   %Analyzed together, closure phase and visiblities are used for imaging reconstruction. %__________________________________________________________________ ",
        "conclusions": "\\label{conclusect} In this paper, we report on high accuracy interferometric measurements of \\object{Betelgeuse} obtained with IOTA/IONIC up to a resolution of $\\sim$11 mas. The precision on squared visibilities enabled by the use of guided optics was better than a percent. The complexity revealed by the data led us to use different methods and models which gave important results. We: \\begin{itemize} \\item{determine the limb-darkened diameter of \\object{Betelgeuse} using parametric and MARCS models as well as a new effective temperature,}  \\item{detect a fully resolved circumstellar environment as bright as 4\\% of the total flux. It is very likely that the silicate dust shell is a significant contributor to this environment, and}  \\item{image two spots at the surface of the star. One of the spots is unresolved and the other one is as large as one half of the stellar radius.} \\end{itemize}    %    simu_john..  %Ouverture sur travaux Andrea. No MOLsphere has been detected in agreement with previous work at near-infrared wavelengths. We thus find a compatibility with the predictions published in \\cite{1975ApJ...195..137S} about the existence of large structures on the surface of \\object{Betelgeuse} due to convection. As \\cite{2009arXiv0902.2602K} observed for other red supergiants, we note that these structures are responsible for up to $\\sim$10 percent of the total flux.  If one assumes that \\object{Betelgeuse} and the spot T1 radiate like blackbodies,  an effective temperature of 3600 K for \\object{Betelgeuse} leads to a spot T1 effective temperature of 4125 K (11 mas diameter and a flux ratio of 0.085). This is compatible with \\cite{1975ApJ...195..137S} who predicted up to 1000 K temperature difference between the hottest and coolest parts of a convective surface. If these spotty structures are similar to those observed on the Sun, then they could be faculae which are bright zones surrounding dark convection cells. Their presence on \\object{Betelgeuse} would show that a magnetic field plays an important role as it is the case for the Sun.        %\\citep{2003A&A...399L...1K}.  %\\bibitem[Krivova et al.(2003)]{2003A&A...399L...1K} Krivova, N.~A., Solanki, S.~K., Fligge, M., \\& Unruh, Y.~C.\\ 2003, \\aap, 399, L1    To go further, it seems that the analysis of this dataset could be improved by injecting small scale ($<$ 10 mas) structures in the modeling. This can be done by comparing our data to simulated visibilities and closure phases generated from RSG hydro-radiative simulations \\citep[such as presented in][]{2009arXiv0907.1860C}. This would undoubtedly bring new clues for the convective nature of the structures found in the present work. The convective motions could play an important role in the mass-loss process operating in RSGs, because of the strong decrease in effective gravity due to turbulent pressure \\citep{2007A&A...469..671J}. The bright plume found by \\cite{2009arXiv0907.1843K} might be the first direct visible evidence of a mass loss mechanism in RSGs. Its formation could be explained by large upwards moving convective cell or by rotation, through the presence of a hot spot at the location of the polar cap of the star. The possibility that this plume is connected to a spot on \\object{Betelgeuse} should be investigated with interferometric techniques in the near future. These perspectives hold the promise of determining a first estimation of the convection scale in RSGs which would represent a new step in the understanding of  the mass loss enigma in RSGs.   %The surface of \\object{Betelgeuse} can only be described by structure that are possibly convective cells , %Josselin 2007 :La convection implique onde de chocs et mouvement supersoniques, dans l'amosph\u008fre, ca peut entra\u0094ner la perte de masse %Echelle de la convection important pour contraindre les mod\u008fles qui peuvent v\\'{e}rifier les taux de perte de masse correspondant....  %proposed model: imagination cr\\'{e}ativit\\'{e}, besoin d'observation, perspective , id\\'{e}e un peu folle , je me l\u0089che, IMAGINATION %s engager    %lien entre dirty beam et figure 9 ; int\\'{e}ressant. % relecture,corrections %Tout est il utile dans la biblio ?"
    },
    "0910/0910.1780_arXiv.txt": {
        "abstract": "We study the generation of a stochastic gravitational wave (GW) background produced by a population of neutron stars (NSs) which go over a hadron-quark phase transition in its inner shells. We obtain, for example, that the NS phase transition, in cold dark matter scenarios, could generate a stochastic GW background with a maximum amplitude of $h_{\\rm BG} \\sim 10^{-24}$, in the frequency band $\\simeq 20-2000\\, {\\rm Hz}$ for stars forming at redshifts of up to $z\\simeq 20.$ We study the possibility of detection of this isotropic GW background by correlating signals of a pair of `advanced' LIGO observatories. ",
        "introduction": "It is widely accepted that neutron stars (NSs) are born with fast rotating angular velocities. Due to magnetic torques, however, the NS periods could well be spun down. This spin-down causes a reduction in the centrifuge force and, consequently, the central energy density of the NSs increases. Those stars, born with densities close to that of the quark deconfinement, may undergo a phase transition forming a strange quark matter core. It is worth stressing, however, that it is not a common sense that such a transition really takes place.  As a consequence of such a putative phase transition, the stars could suffer a collapse which could excite mechanical oscillations. The great amount of energy generated in this process, $\\Delta E\\sim 10^{53}\\,ergs$ ($\\sim 0.1\\, M_{\\odot}c^{2}$) \\cite{marr1}, could be dissipated, at least partially, in the form of gravitational waves (GWs), which are studied in the present work.  We adopted in the present paper the history of star formation derived by Springel \\& Hernquist \\cite{springel}, who employed hydrodynamic simulations of structure formation in a $\\Lambda$ cold dark matter ($\\Lambda$CDM) cosmology. These authors study the history of cosmic star formation from the ``dark ages\", at redshift  ${\\rm z} \\sim 20$, to the present.  Besides the reliable history of star formation by Springel \\& Hernquist, we consider the role of the parameters $\\alpha_{i=1,2}$, which gives the fraction of the progenitor mass which forms the remnant NSs. We consider that the remnant mass is given as a function of the progenitor mass, namely $M_r=\\alpha_1m+\\alpha_2$. Later on we justify why $M_r$ is written in this way. Recall that a given initial mass function (IMF) refers to the distribution function of the stellar progenitor mass, and also to the masses of the remnant compact objects left as a result of the stellar evolution.  In the present study we have adopted a stellar generation with a Salpeter IMF, which is consistent with Springel \\& Hernquist, since they show that population II stars could have been formed at high redshift too. We then discuss briefly what conclusions would be drawn whether (or not) the stochastic background studied here is detected by GW observatories such as LIGO and VIRGO.  The paper is organized as follows. Section 2 deals with the NS EOSs adopted in our study; in Section 3 we consider the IMF and the NS masses; Section 4 deals with the GW production; in Section 5 we present the numerical results and discussions; Section 6 the detectability of the background of GWs are considered and finally in Section 7 we present the conclusions. ",
        "conclusions": "We present here a study concerning the generation of GWs produced from a cosmological population of NSs. These stars may undergo a phase transition if born close to the transition density, suffering a micro-collapse and exciting quasi-normal modes.  We show that a detectable background is possible only if the SFR density is much greater than that predicted by Springel \\& Hernquist or if the energy generated in the phase transition is almost completely channeled to excite the f- mode. Obviously, a too optimistic combination of these two possibilities could do the same.  \\ack{JCNA would like to thank CNPq and Fapesp for financial support.}"
    },
    "0910/0910.4393_arXiv.txt": {
        "abstract": "We present a comparison between Gaussian processes (GPs) and artificial neural networks (ANNs) as methods for determining photometric redshifts for galaxies, given training set data.  In particular, we compare their degradation in performance as the training set size is degraded in ways which might be caused by the observational limitations of spectroscopy. Using publicly-available regression codes, we find that performance with large, complete training sets is very similar, although the ANN achieves slightly smaller root mean square errors. Training sets with brighter magnitude limits than the test data do not strongly affect the performance of either algorithm, until the limits are so severe that they remove almost all of the high-redshift training objects. Similarly, the introduction of a plausible number (up to 10\\%) of inaccurate redshifts into the training set has little effect on either method. However, if the size of the training set is reduced by random sampling, the RMS errors of both methods increase, but they do so to a lesser extent and in a much smoother manner for the case of GP regression; for the example presented {\\sc{ANNz}} has RMS errors $\\sim20\\,$\\% worse than GP regression in the small training-set limit. Also, when training objects are removed at redshifts $1.3<z<1.7$, to simulate the effects of the ``redshift desert'' of optical spectroscopy, the Gaussian process regression is successful at interpolating across the redshift gap, while the ANN suffers from strong bias for test objects in this redshift range. Overall, GP regression has attractive properties for photometric redshift estimation, particularly for deep, high-redshift surveys where it is difficult to obtain a large, complete training set. At present, unlike the ANN code, public GP regression codes do not take account of inhomogeneous measurement errors on the photometric data, and thus cannot estimate reliable uncertainties on the predicted redshifts.  However, a better treatment of errors is in principle possible, and the promising results in this paper suggest that such improved GP algorithms should be pursued. ",
        "introduction": "\\subsection{Photometric redshift estimation} The idea of estimating galaxy redshifts from broad-band photometry was first applied by \\citet{Baum62} to obtain redshifts for suspected members of a cluster which were too faint for spectroscopy using the instruments of the time. Baum determined the redshifts by first measuring photometry of a set ellipticals in known redshift clusters, to create a coarsely-sampled template spectral energy distribution, and then shifting this in logarithmic wavelength space (as if redshifted) until it matched the photometry of the unknown objects.   Since then it has become possible to fit to a wide variety of higher resolution templates, based on empirical \\citep[e.g.][]{cww} or synthetic \\citep[e.g.][]{bc03} galaxy spectra.  Model photometry is obtained by directly multiplying a redshifted template spectrum by the measured response curves for each photometric band, and is compared with data to determine the best-fitting redshift.  Template-fitting methods are very widely used \\citep[e.g.][]{hyperz,bpz}, since they can be applied to any photometric dataset provided that the instrument response curves are known and the templates appropriate for the galaxy being studied.  However, there are often systematic uncertainties in both of these; filter curves, detector responses and atmospheric transmission are in general not precisely known, and spectral templates are necessarily composed of or calibrated against low-redshift galaxies, and become less reliable as redshift increases.  Some of these problems with template fitting can be mitigated by using known redshifts for calibration \\citep[e.g.][]{ilbert06,feldmann2006}. However, where a representative subset of the objects have precisely known redshifts, it is much simpler to use these objects to directly train some function so that it returns redshift as a function of photometric data, and then apply this function to the remainder of the catalogue.  Empirical training-set methods of this kind have the additional advantage that it is straightforward to include information from additional catalogued parameters such as angular size or surface brightness profile.  Empirical photometric redshift methods are currently dominated by artificial neural networks \\citep[ANNs;][]{firth2003,Vanzella2004}. ANNs can perform general non-linear mappings (e.g. from photometric data to redshift) by the use of multiple layers of linear combinations of data (optionally with non-linear transfer functions on some of the nodes), with the weights on these combinations being free parameters which are modified during training.  Simple feed-forward networks (in the sense that values propagate in one direction only, e.g. from photometry, via intermediate layers of linear combinations, to the estimated redshift) have the useful property, during training, that the errors on the redshift estimates help inform the direction in which to modify the weights, and they can thus be trained quite efficiently.  The most commonly-used implementation of ANNs for redshift estimation is the publicly-available {\\sc{ANNz}} package \\citep{annz04}.  \\citet{megazlrg} show that {\\sc{ANNz}}  matches or surpasses the performance of the best publicly-available template-fitting codes, when it is provided with a high quality training set.  However, neural networks are far from the only tools able to perform non-linear regression. One alternative, which has recently been favoured in the machine-learning community, is a class of models known as Gaussian processes.   \\subsection{Gaussian process regression} A Gaussian process (GP) is defined simply as a collection of random variables which have a joint Gaussian distribution.  The utility of GPs for regression is that they can be used as a prior distribution over a space of functions which are considered as possible models for the data \\citep[e.g.][]{wr96}.  The space of functions represented by a Gaussian process is completely defined by a mean function (usually taken to be zero) and a covariance function, which for $n$ training data ${\\bf{x}}_1...{\\bf{x}}_n$ (each of which would be a vector of photometry in our case) and one test point ${\\bf{x}}_*$ (i.e. the photometry of a galaxy with unknown redshift) takes the form of an $(n+1) \\times (n+1)$ matrix of covariances $k_{ij}=k({\\bf{x}}_i,{\\bf{x}}_j)$ between the data.  Thus, the vector of ``outputs'', which contains redshifts for the training data $y_1...y_n$ and the redshift to be inferred for the test point $y_*$, is assumed to be given by:  \\begin{equation} \\label{eq:smalldef} \\left( \\begin{array}{c} y_1 \\\\ \\vdots \\\\ y_n \\\\ y_* \\end{array} \\right) = \\mathcal{N}\\left({\\bf{0}}, \\left[ \\begin{array}{cccc} k_{11} + \\sigma_n^2& \\cdots & k_{1n} & k_{1*} \\\\ \\vdots & \\ddots & & \\vdots \\\\ k_{n1} &  &  k_{nn} +\\sigma_n^2 & k_{n*} \\\\ k_{*1} & \\cdots &  k_{*n} & k_{**} \\end{array} \\right] \\right) \\end{equation}  where $\\mathcal{N}({\\bf{0}},C)$ is a joint Gaussian distribution with mean ${\\bf{0}}$ and covariance matrix $C$, and $\\sigma_n^2$ is a noise term added to the first $n$ diagonal elements, both to account for noise in the data (assumed independent and identically distributed) and to increase numerical stability.  While this distribution appears extremely simple, it should be noted that there is considerable freedom to choose the covariance function $k({\\bf{x}}_i,{\\bf{x}}_j)$, also commonly referred to as the kernel, which can be any positive definite function.  This means we can achieve non-linear regression by using a non-linear kernel. We will discuss our specific choice of kernel later, but in general one chooses some function whose value indicates the similarity of datapoints ${\\bf{x}}_i$ and ${\\bf{x}}_j$, and which may have a number of hyperparameters whose optimisation forms part of the training process.  If we define \\begin{displaymath} K = \\left( \\begin{array}{ccc} k_{11} & \\cdots & k_{1n} \\\\ \\vdots & \\ddots &  \\\\ k_{n1} &  &  k_{nn} \\end{array} \\right) , \\,{\\bf{k}}_* = \\left( \\begin{array}{c} k_{1*} \\\\ \\vdots \\\\ k_{n*} \\end{array} \\right), {\\bf{y}} = \\left( \\begin{array}{c} y_{1} \\\\ \\vdots \\\\ y_{n} \\end{array} \\right) \\end{displaymath} then it can be shown \\citep[e.g.][]{rw06} that the posterior probability distribution of the predicted redshift $y_*$ is a Gaussian distribution, with mean $\\hat{y}_*$ and variance $\\sigma_*^2$ given by \\begin{eqnarray}\\label{eq:gprmean} \\hat{y}_* & = & {\\bf{k}}_*^T(K+\\sigma_n^2I)^{-1}{\\bf{y}} \\\\ \\label{eq:gprvar} \\sigma_*^2 & = & k_{**} - {\\bf{k}}_*^T(K+\\sigma_n^2I)^{-1}{\\bf{k}}_* \\end{eqnarray} where $A^T$ denotes the transpose of $A$, $A^{-1}$ denotes the inverse of $A$, and $I$ is the identity matrix.   \\subsection{Previous work}  \\citet{Way06} first evaluated Gaussian processes for photometric redshift estimation, and found that an ensemble of neural networks produced a slightly smaller RMS error.  However, they used a smaller training set for the GP-based method due to computational limitations (the na\\\"{\\i}ve implementation of equations \\ref{eq:gprmean} and \\ref{eq:gprvar} requires the inversion of an $n\\times n$ matrix, which involves $O(n^3)$ operations) so this was not a completely fair comparison of the methods.  More recently, \\citet{foster09} have shown that it is possible to use rank-reduction techniques which allow inference using large training sets, with $n$ objects, to be performed in $O(nm^2)$ operations (where $m<n$). Such methods make use of a full covariance matrix for a subset of only $m$ objects (an $m\\times m$ matrix), in conjunction with an $m\\times n$ matrix of covariances between the full and partial training sets.  With the ability to use information from a large training set, of the same size as used by ANN-based methods, \\citet{Way09} find that GP regression yields a slightly lower RMS error than {\\sc{ANNz}}, when applied to galaxies drawn from the Sloan Digital Sky Survey \\citep[SDSS;][]{sdss2000}.  Way et al.~also characterise the performance of GP regression using different kernel functions and choices for the size of the reduced rank $m$.  These existing evaluations of GP regression for photometric redshift estimation have used an SDSS-type survey as a model.  That is, they have assumed that a very large set of training objects is available, and have implicitly ensured that the training set is absolutely complete by selecting training and test objects in the same way from the SDSS spectroscopic sample (i.e. those SDSS galaxies which have known spectroscopic redshifts).  Use of SDSS data has also resulted in the exploration of a rather narrow range in redshift.  In this paper, we extend this work by considering a number of observationally-motivated restrictions on the training sets.  In particular, we consider cases where the number of available spectra are small, but complete, and cases in which the training set is limited in magnitude in a different manner than the test data.  We also examine the effect of reducing the number training objects with spectral types lacking emission lines, or in the emission line ``desert'' at redshifts $1.3<z<1.7$, and compare the resistance of the methods to ``bad'' training objects with incorrectly determined redshifts.  We use photometry simulated to represent a deep, high-redshift survey. ",
        "conclusions": "We have compared a simple implementation of GP regression, based on the {\\sc{stableGP}} code of \\citet{foster09}, with the popular neural network code {\\sc{ANNz}}.  With large, complete training sets the precision of the methods is very similar, with {\\sc{ANNz}} achieving slightly smaller RMS errors, particularly at redshifts $z<1$. However,   we note that GP regression, unlike {\\sc{ANNz}}, does not require an additional set of validation data to preserve generality, and so could take advantage of a larger training sample.  We investigate placing magnitude limits on the training sets but find that the performance of neither algorithm is strongly affected, until the limits are so severe that they remove a large fraction of the training objects in the redshift range of the test object.  In this case the GP algorithm produces less biased results but with a larger scatter.  Similarly, the introduction of a plausible number (up to 10\\%) of inaccurate redshifts into the training set, or the removal of objects of a particular spectral type, has little effect on either method.  If the size of training set is instead reduced by random sampling, the RMS errors of both methods increase, but they do so to a lesser extent and in a much smoother manner for the case of GP regression.  With this particular dataset and network architecture we find that the ANN method deteriorates sharply with training sets smaller than about 2000 objects, where its RMS errors are $\\sim20$\\% larger than for GP regression, with considerably increased bias.  While the GP results could in principle be reproduced by some appropriately selected ANN architecture, the computational difficulty of selecting and then training such a network may make the GP method preferable.  In addition, when training objects are removed at redshifts $1.3<z<1.7$, to simulate the effects of the ``redshift desert'' of optical spectroscopy, the GP regression is successful at interpolating across the redshift gap, while {\\sc{ANNz}} suffers from strong bias for test objects in this redshift range.    These properties  make GP regression a very attractive algorithm for the estimation of photometric redshifts, particularly in the case where few training redshifts are available at the redshift of interest.  Since telescopes of any given aperture are in general able to image fainter objects than they can obtain spectra for, this is likely to be the situation for the vast majority of both present and future deep, high-redshift surveys.  At present, unlike the ANN code, public GP regression codes do not take account of inhomogeneous measurement errors on the photometric data, and thus cannot estimate reliable uncertainties on the predicted redshifts.  Such uncertainty values can be used to form subsamples of objects with, on average, more accurate redshift estimates.  While such subsamples have limited applications due to the complex selection effects involved, redshift uncertainty values  can also  be used (e.g. using a Monte Carlo method) to calculate more robust statistics about the whole population of galaxies.  The advantages of GPs demonstrated in this paper suggest that such improved GP algorithms should be pursued, and we are exploring both Monte Carlo and analytic options for improving their treatment of errors."
    },
    "0910/0910.1864_arXiv.txt": {
        "abstract": "{Bolometric Interferometry is a technology currently under development that will be first dedicated to the detection of B-mode polarization fluctuations in the Cosmic Microwave Background. A bolometric interferometer will have to take advantage of the wide spectral detection band of its bolometers in order to be competitive with imaging experiments. A crucial concern is that interferometers are presumed to be importantly affected by a spoiling effect known as bandwidth smearing.} {In this paper, we investigate how the bandwidth modifies the work principle of a bolometric interferometer and how it affects its sensitivity to the CMB angular power spectra.} {We obtain analytical expressions for the broadband visibilities measured by broadband heterodyne and bolometric interferometers.  We investigate how the visibilities must be reconstructed in a broadband bolometric interferometer and show that this critically depends on hardware properties of the modulation phase shifters. If the phase shifters produce shifts that are constant with respect to frequency, the instrument works exactly as a monochromatic one (the modulation matrix is not modified), while if they vary (linearly or otherwise) with respect to frequency, one has to perform a special reconstruction scheme, which allows the visibilities to be reconstructed in frequency sub-bands. Using an angular power spectrum estimator accounting for the bandwidth, we finally calculate the sensitivity of a broadband bolometric interferometer. A numerical simulation has been performed and confirms the analytical results.} {We conclude (i) that broadband bolometric interferometers allow broadband visibilities to be reconstructed whatever the kind of phase shifters used and (ii) that for dedicated B-mode bolometric interferometers, the sensitivity loss due to bandwidth smearing is quite acceptable, even for wideband instruments (a factor 2 loss for a typical 20\\% bandwidth experiment).} {} ",
        "introduction": " ",
        "conclusions": "In this article we have analytically and numerically studied the work principle of a broadband bolometric interferometer. We have defined its (indirect) observables -- the broadband visibilities -- and introduced numerical methods to reconstruct them. We have finally calculated the sensitivity of such an instrument dedicated to the B-mode.  Bolometers are naturally broadband, and consequently the design of a broadband bolometric interferometer is identical to the design of the monochromatic one described in [C08]: a broadband bolometric interferometer only needs broadband components (horns, filters, etc.). Nevertheless, we have seen that the modulation matrix that should be used to reconstruct the broadband visibilities depends on some hardware properties of the modulation phase shifters. If these are constant with respect to frequency, the modulation matrix should be the one defined in [C08] for a monochromatic instrument. If they are dependent on frequency (this dependence should be known of course), a more complicated scheme involving a reconstruction in sub-bands of the intensity visibilities should be performed. We have checked with a numerical simulation that in both cases the visibilities can be reconstructed without other loss in sensitivity than the one due to the smearing.  Visibilities are defined as the convolution of the Fourier transform of the signal with a kernel, which is defined as the Fourier transform of the primary beam in the monochromatic case. We have shown that the effect of the smearing is to stretch this kernel, in the baseline direction only, and that the amplitude of the smearing only depends on three quantities: the bandwidth, the baseline length and the size of the primary beam. We have finally defined, as a function of broadband visibilities, a new power spectrum estimator and derived from it a generalized uncertainty formula.  The main conclusion of this article is that for a bolometric interferometer dedicated to CMB B-mode, the sensitivity loss -due to bandwidth smearing- is quite acceptable (a factor 2 loss for a typical 20\\% bandwidth experiment)."
    },
    "0910/0910.4736_arXiv.txt": {
        "abstract": "The dark energy component of the cosmic budget is represented by a self-interacting scalar field. The violation of the null energy condition is allowed. Hence, such component can also represent a phantom fluid. The model is tested using supernova type Ia and matter power spectrum data. The supernova test leads to preferred values for configurations representing the phantom fluid. The matter power spectrum constraints for the dark energy equation of state parameter are highly degenerated. In both cases, values for the equation of state parameter corresponding to the phantom fluid are highly admitted if no particular prior is used. \\newline \\leftline{Pacs: 98.80.-k,95.36.+x} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.3927_arXiv.txt": {
        "abstract": "The phase transition from hadronic to quark matter may take place already during the early post-bounce stage of core collapse supernovae when matter is still hot and lepton rich. If the phase transition is of first order and exhibits a barrier, the formation of the new phase occurs via the nucleation of droplets. We investigate the thermal nucleation of a quark phase in supernova matter and calculate its rate for a wide range of physical parameters. We show that the formation of the first droplet of a quark phase might be very fast and therefore the phase transition to quark matter could play an important role in the mechanism and dynamics of supernova explosions. ",
        "introduction": "The possibility that a first-order phase transition in dense matter has implications for the explosions of supernovae was first proposed by Migdal {\\it et al.} about 30 years ago \\cite{1979PhLB...83..158M}. Since then, a large number of papers addressed this issue with different hypotheses on the nature of the phase transition: pion or kaon condensate matter or quark matter, different models to compute the equation of state (EoS), different simplifications for the complex hydrodynamical evolution of collapsing massive stars and different approaches for the neutrino Boltzmann transport \\cite{1986Ap&SS.119...45T,1985PhRvL..55..126B,Gentile:1993ma}. However, only very recently was it possible to perform simulations using general relativistic Boltzmann neutrino transport equations and adopting realistic equations of state for quark matter \\cite{Nakazato:2008su,Sagert:2008ka}. In Ref.~\\cite{Sagert:2008ka}, in particular, it was shown that the phase transition to quark matter can occur already during the early post-bounce phase of a core collapse supernova event and that it produces a second shock wave (the first being the usual shock wave after the bounce) which triggers a delayed supernova explosion, even within spherical symmetry, for masses of the progenitor star up to 15 $M_{\\odot}$. The formation of quark matter is also responsible for the emission of a neutrino burst, typically a few hundred milliseconds after the first neutronization burst, which could be detected by currently available neutrino detectors, representing a spectacular possible signature of quark matter formation in compact stars (see also \\cite{Benvenuto:1989qr}).  In the studies mentioned above, an important physical phenomenon is neglected for the sake of simplicity: the process of phase conversion in a first-order transition is actually driven by the nucleation of finite-size structures, such as droplets or bubbles, of the new phase within the old phase. The surface tension, $\\sigma$, is the physical quantity that determines the nature of the process of phase conversion. If $\\sigma$ is sufficiently small, nucleation can be very fast and the new phase is produced almost in mechanical, thermal and chemical equilibrium with the nuclear matter phase. An intermediate value of $\\sigma$ might render nucleation very difficult and the nuclear phase can be metastable for a significant amount of time. Then, the process of formation of the new phase, once triggered, would be a genuine nonequilibrium process, in which different mechanisms can take place: deflagration, detonation, convective instabilities (see, e.g., Refs. \\cite{Drago:2005yj,Drago:2008tb} and references therein). Finally, nucleation is highly suppressed for large values of $\\sigma$, and the formation of the new phase may only proceed via spinodal decomposition, if the density achieved is high enough to flatten out the activation barrier. (See Ref.~\\cite{Bessa:2008nw} for a recent detailed discussion of the phase conversion process in a first-order phase transition.)  The nucleation of quark matter in neutron stars has been explored mainly within a scenario, proposed in Ref.~\\cite{Pons:2001ar}, in which the formation of quark matter occurs only when the protoneutron star is almost completely deleptonized and the temperature has already dropped to, say, $1$ MeV. Under these conditions, quantum nucleation has been shown to be the most important mechanism for the formation of a quark phase \\cite{Iida:1997ay,Iida:1998pi,Berezhiani:2002ks,Bombaci:2004mt,Bombaci:2008wg}, also when color superconducting quark phases are present \\cite{Drago:2004vu,Bombaci:2006cs}. A less explored scenario is the nucleation of quark matter in hot and lepton-rich protoneutron stars. Refs.~\\cite{Horvath:1992wq,Olesen:1993ek} present the first estimates and calculations showing thermal nucleation to be very efficient for temperatures of roughly $10$ MeV and practically negligible for temperatures below $2$ MeV. The simplifying assumption adopted in those papers (no leptons are included in the quark equation of state) might be, however, not realistic considering that for the large temperatures needed to nucleate quark matter the neutrino mean free path is small, and therefore neutrinos are trapped. In the presence of neutrinos, the critical densities for the phase transition are shifted toward larger values compared to a deleptonized neutron star. Moreover, as we will discuss in the following, the assumption of flavor conservation during the phase transition (also adopted in Ref.~\\cite{Lugones:1997gg}) might be too conservative in light of the quark density fluctuations that are evidently present.  The main goal of this paper is to compare the time scale associated with the phase conversion with the dynamical time scale of core collapse supernovae during which quark matter might be eventually formed. In particular, since the explosion process occurs within a time scale of few hundred ms, the nucleation time must be of the same order of magnitude if quark matter plays indeed a role in the explosion mechanism.  For this purpose, we reconsider thermal nucleation of quark matter within the scenario proposed in Ref.~\\cite{Sagert:2008ka} of a phase transition occurring already during the early post-bounce stage of a core collapse supernova. This implies a value for the critical density $n_{c}\\lesssim 2 n_0$ (where $n_0=0.16$ fm$^{-3}$ is the nuclear matter saturation density).  We systematically investigate the windows of free parameters of the model adopted to compute the equations of state and the corresponding nucleation time scale. Our strategy is to present {\\it underestimates} of this time scale, so that a successful phase conversion could somehow constrain the equation of state parameter space. We also discuss the important issue of flavor number conservation during the phase transition. We argue that thermal nucleation might indeed be efficient under the conditions realized in a star soon after bounce. In such case, the phase transition would proceed very fast, thus confirming the results found in Ref.~\\cite{Sagert:2008ka}.  The paper is organized as follows. In Sec. II we present the model we adopt for the equation of state, as well as a brief self-contained description of homogeneous nucleation and the role of statistical fluctuations. In Sec. III we present our results for the nucleation of nonstrange and strange quark matter. There, we also discuss quantum nucleation and the spinodal instability. Besides, we verify our equations of state by computing the stellar structure that emerge form the Tolman-Oppenheimer-Volkov (TOV) equations. Finally, Sec. IV contains our conclusions. ",
        "conclusions": "We investigated the possibility of the formation of quark matter in supernova matter, i.e. for temperatures of the order of a few tens of MeV and in the presence of trapped neutrinos, assuming that the corresponding critical density does not exceed $2 n_0$.  We argued that thermal nucleation of quark phase droplets is eventually the dominant mechanism for the formation of the new phase and that, due to fluctuations in the number densities of quarks, the phase transition involves directly the beta equilibrated quark phase. We have calculated the nucleation rate for different values of the free parameters. The surface tension, as expected, is the physical quantity which mainly controls the nucleation process.  Among the different equations of state and conditions at bounce we have tested, the choice $T=20$ MeV and $Y_L=0.4$ is the most likely to occur. Within this choice of physical conditions, if the phase transition involves only noninteracting up and down quarks, a value of $\\sigma$ smaller than $\\sim15~$MeV/fm$^2$ should be required for nucleation to be efficient. Such a low value of $\\sigma$ is compatible with lattice QCD calculations, although they are obtained at large temperatures and small densities \\cite{deForcrand:2006ec}. On the other hand, phenomenological estimates at large density and zero temperature indicate larger values of $\\sigma$ \\cite{Voskresensky:2002hu}. Therefore, we consider this scenario to be unlikely.  If strange quarks are produced during the phase transition, we conclude that, if $\\sigma$ is smaller than $\\sim 10~$MeV/fm$^2$ (for $m_s=100$MeV and $c=0.3$), the nucleation time for the first drop of quark matter is sufficiently small and the appearance of quark matter can indeed strongly affect the supernova evolution as shown in \\cite{Sagert:2008ka}. The cold hybrid stars obtained after deleptonization and cooling have, if perturbative QCD corrections are included in the equation of state, maximum masses compatible with recent pulsar observations.   {\\it Note added}- After finishing this work a paper has been published which discusses similar issues \\cite{Bombaci:2009jt}. In that work, the authors also conclude that nucleation is possible in protoneutron star matter. The main difference between our work and Ref.~\\cite{Bombaci:2009jt} is that in the latter the thermal nucleation of quark matter is studied within hot and deleptonized hadronic matter."
    },
    "0910/0910.2274_arXiv.txt": {
        "abstract": "We develop the finite temperature theory of $p$-adic string models. We find that the thermal properties of these non-local field theories can be interpreted either as contributions of standard thermal modes with energies proportional to the temperature, or inverse thermal modes with energies proportional to the inverse of the temperature, leading to a {\\it thermal duality} at leading order (genus one) analogous to the well known T-duality of string theory. The $p$-adic strings also recover the asymptotic limits (high and low temperature) for arbitrary genus that purely stringy calculations have yielded. We also discuss our findings surrounding the nature of the Hagedorn transition. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.3595.txt": {
        "abstract": "{ Deeply inside dense molecular clouds and protostellar disks, the interstellar ices are protected from stellar energetic UV photons. However, X-rays and energetic cosmic rays can penetrate inside these regions triggering chemical reactions, molecular dissociation and evaporation processes. We present experimental studies on the interaction of heavy, highly charged and energetic ions (46 MeV $^{58}$Ni$^{13+}$) with ammonia-containing ices H$_2$O:NH$_3$ (1:0.5) and H$_2$O:NH$_3$:CO (1:0.6:0.4) in an attempt to simulate the physical chemistry induced by heavy ion cosmic rays inside dense astrophysical environments. The measurements were performed inside a high vacuum chamber coupled to the IRRSUD (IR radiation SUD) beamline at the heavy ion accelerator GANIL (Grand Accelerateur National d'Ions Lourds) in Caen, France. The gas samples were deposited onto a polished CsI substrate previously cooled to 13 K. \\textit{In-situ} analysis is performed by a Fourier transform infrared spectrometer (FTIR) at different fluences. The averaged values for the dissociation cross section of water, ammonia and carbon monoxide due to heavy cosmic ray ion analogs are $\\sim 2 \\times$10$^{-13}$, 1.4$\\times$10$^{-13}$ and 1.9$\\times$10$^{-13}$ cm$^2$, respectively. In the presence of a typical heavy cosmic ray field, the estimated half life for the studied species is 2-3 $\\times 10^6$ years. The ice compaction (micropore collapse) due to heavy cosmic rays seems to be at least 3 orders of magnitude higher than the one promoted by (0.8 MeV) protons . The infrared spectra of the irradiated ice samples present lines of several new species including HNCO, N$_2$O, OCN$^-$ and NH$_4^+$ . In the case of the irradiated H$_2$O:NH$_3$:CO ice, the infrared spectrum at room temperature reveals five bands that were tentatively assigned to vibration modes of the zwitterionic glycine (NH$_3^+$CH$_2$COO$^-$). ",
        "introduction": "The birthplace of stars is the densest and coldest place of the interstellar medium (ISM) called dense molecular clouds. These regions have typical gas densities of 10$^3$-10$^8$ atoms cm$^{-3}$ and temperatures of the order of 10-50 K. Due to the low temperature, dust particles (mainly silicates and carbonaceous compounds like amorphous C and SiC) can accrete molecules from the gas phase and get coated with an ice mantle. The interstellar ice mantles (or simply interstellar ices) are composed primarily of amorphous H$_2$O but usually also contain a variety of other simple molecules such as CO$_2$, CO, CH$_3$OH and NH$_3$ (e.g. Ehrenfreund \\& Charnley 2000; Ehrenfreund \\& Shuttle 2000; Gibb et al. 2001; Boogert et al. 2004). Laboratory studies and astronomical observations indicate that photolysis and radiolysis of such ices can create complex organic compounds, and even pre-biotic molecules like amino acids and nucleobases (e.g. Bernstein et al. 2002; Mu\u00f1oz Caro et al. 2002; Kobayashi et al. 2008).  The observation of molecules in the gas phase deeply inside these cold and dense regions, where the gas sticking efficiency on grains is close to unity, suggests that they are indeed energetically active regions (e.g. Ehrenfreund \\& Charnley 2000). Only a negligible amount of stellar UV photons reaches the inner parts of dense regions (Tielens \\& Hagen 1982; d'Hendecourt et al. 1985). Other mechanisms have therefore been proposed to explain the presence of free molecules such as cosmic-ray-induced UV photons (Prasad \\& Tarafdar 1983) and direct cosmic ray interaction with ice mantles (e.g Shen et al. 2004). Both processes lead to molecular desorption from the surface. As pointed out by Shen et al. (2004), in the case of cosmic ray particles, the sputtering promoted by direct impact and the whole grain heating (classical sublimation) due to energy deposition in the bulk are also two efficient processes to release molecules from frozen surfaces to the gas phase.  The first detection of ammonia ices has occurred around the young stellar source (YSO) NGC 7538 IRS9 (Lacy et al. 1998).  It has also been observed around the massive protostar GCS3 (Chiar et al. 2000) and several other massive YSO like W33A (Gibb et al. 2000; 2001). Ammonia has also been detected extensively inside solar system in cometary coma and on several moons (e.g. Kawakita et al. 2006; Bird et al. 1999; Moore et al. 2007 and references therein).  One of the most reactive and important molecules observed from photolysis and radiolysis experiments of ammonia-containing ices, in which there is a carbon source, is the cyanate ion, OCN$^-$ (e.g. Demyk et al. 1998; Hudson \\& Moore 2000; van Broekhuizen et al. 2005 and references therein). Its infrared band around 2165 cm$^{-1}$ was observed in several protostellar sources (e.g. Lacy et al. 1984; Gibb et al. 2000; Whittet et al. 2001) indicating regions were UV photochemistry and ion/electron bombardments have played an important role in the chemical alteration of grain mantles.  Although the flux of heavy ions (e.g. Fe, Ni, Si, Mg, ...) is about 3-4 orders of magnitude lower than that of protons (Roberts et al. 2007; Mennella et al. 2003), their effects play an important role on interstellar grains since they can deposit  about 100 times more energy than the light ions (He$^+$ and protons) inside the grains. Brown et al. (1984) and Fama et al. (2008) using high energy He$^+$ and protons have found that the sputtering yields scales with the square of the electronic stopping power. We have recently proved that this square law extends its validity up to values of electronic stopping, such as those induced by swift heavy ions (Seperuelo Duarte et al. 2009a,b). Consequently, amount of species released per impact to gas phase due to heavy ions could be 4-5 orders of magnitude higher than those for protons.  In order to characterize the effects of heavy ions on different interstellar ice analogs, it is necessary to measure their sputtering yields and chemical reaction cross sections and to compare them to those promoted by He$^+$ and protons. Studies on the effect of heavy projectiles in astrophysical ice analogs are scarce. Most of experiments have been performed in the keV or hundred of keV range, mainly with protons, He and Ar ions (e.g. Loeffler et al. 2006; Gomis et al. 2004a,b). Due to the low kinetic energy of these ions, the structural and chemical changes in the frozen ices are mainly governed by nuclear energy loss in elastic ion-target atom collisions (e.g. Brown et al. 1982; 1984).  In this work, we present infrared measurements of ammonia-containing astrophysical ice analogs (H$_2$O:NH$_3$ (1:0.5) and H$_2$O:NH$_3$:CO (1:0.6:0.4)) irradiated by 46 MeV $^{58}$Ni$^{13+}$ to simulate the modification induced by heavy cosmic rays inside dense astrophysical environments. At such high ion velocity, the energy deposition  is mainly due to inelastic electronic interactions with target electrons (electronic stopping power regime). ",
        "conclusions": "Experimental study on the interaction of heavy, highly charged and energetic ions (46 MeV $^{58}$Ni$^{13+}$) with ammonia-containing ices H$_2$O:NH$_3$ (1:0.5) and H$_2$O:NH$_3$:CO (1:0.6:0.4) is performed to simulate the physical chemistry induced by cosmic rays inside dense regions of interstellar medium like dense molecular clouds or protoplanetary disks. Our conclusions are: \\begin{enumerate} \\item At the beginning of the irradiation, the ices become compact as revealed by the particular dependence of the column density on the fluence. Measuring the degree of compaction  through the integrated area of the OH dangling bond feature ($\\sim$ 3650 cm$^{-1}$) from of water molecules trapped into ice micropores, the effect promoted by the impact of heavy ions seems to be at least 3 orders of magnitude higher than the one promoted by (0.8 MeV) protons. \\item The infrared spectra of the irradiated samples present the production of several new species, including OCN$^-$, HNCO and NH$_4^+$. The OCN$^-$ band (2169 cm$^{-1}$) is observed even at room temperature. During the ice heating a broad feature was observed around 2218-2200 cm$^{-1}$ and was tentatively assigned to non-volatile nitriles. At room temperature an intense and sharp feature was observed at ($\\sim$ 2150 cm$^{-1}$) and was tentatively assigned to aliphatic isocyanide ($\\sim$ 2150 cm$^{-1}$). \\item The spectra of the irradiated  H$_2$O:NH$_3$:CO ice (1:0.6:0.4) at room temperature reveal several bands which are tentatively assigned to zwitterionic glycine and HMT. Despite the presence of these complex molecules in the organic residue from the bombardments of interstellar ice analogs by energetic and heavy ions has to be confirmed. \\item The obtained value for the dissociation cross section of water, ammonia and carbon monoxide (in the tertiary mixture ice) due to heavy cosmic ray analogs are $\\sim 2 \\times 10^{-13}$, 1.4$\\times 10^{-13}$ and 1.9$\\times 10^{-13}$ cm$^{2}$, respectively. These values seem not to be affected by small ($<$30\\%) changes in the initial relative molecular abundances. \\item In the presence of a typical heavy cosmic ray field the estimated half lives for the studied species is about 2-3 $\\times 10^6$ years. This value is in a good agreement with the half lives of typical molecular clouds and protostellar clouds. \\end{enumerate}  Although they represent only a small fraction ($\\sim$ 1\\%) of the cosmic rays flux, some effects (e.g. molecular sputtering and ice compaction) promoted by heavy ions on interstellar ice grains are more intense than those promoted by protons. This should be taken under consideration in future chemical models of inner regions of molecular clouds since heavy ions will contribute molecules to the gas phase and trigger chemical surface and bulk reactions.  % -------Acknowledgements---------------------------- %"
    },
    "0910/0910.2873_arXiv.txt": {
        "abstract": "Inclined air showers -- those arriving at ground with zenith angle with respect to the vertical $\\theta > 60^{\\circ}$ -- are characterised by the dominance of the muonic component at ground which is accompanied by an electromagnetic ``halo\" produced mainly by muon decay and muon interactions. By means of Monte Carlo simulations we give a full characterisation of the particle densities at ground in ultra-high energy inclined showers as a function of primary energy and mass composition, as well as for different hadronic models assumed in the simulations. We also investigate the effect of intrinsic shower-to-shower fluctuations in the particle densities. ",
        "introduction": "Inclined air showers are conventionally defined as those arriving at ground with zenith angles $\\theta$ above $60^\\circ$.  At large zenith angles the electromagnetic~(EM) component in air showers, mainly produced by the decay of $\\pi^0$s, is largely absorbed in the greatly enhanced atmospheric depth the shower needs to cross before reaching ground, so that in a first approximation only the more penetrating particles such as muons survive to ground.  Muons are accompanied by an electromagnetic component produced mainly by muon decay in flight and muon interactions such as bremsstrahlung, pair production and nuclear interactions. A full characterisation of this so-called electromagnetic ``halo\" is given in this work.  The study of inclined showers is of great interest because their detection immediately enhances the exposure of existing air shower detectors by up to about $30\\%$ with respect to that achieved with vertical showers $(\\theta<60^\\circ)$, extending the field of view to sky directions otherwise inaccessible.  Inclined showers have been detected in the past in arrays of ground detectors such as the Haverah Park~\\cite{Ave:1999cp} and AGASA~\\cite{Yoshida:2001pw} experiments. Modern detectors such as the Pierre Auger Observatory~\\cite{Abraham:2004dt} or the Telescope Array~\\cite{TA,Kawai:2008zza} can also be used to detect showers with large zenith angles.  In particular, the Surface Detector Array (SD) of the Pierre Auger Observatory is well suited to detect very inclined showers at energies above about $5 \\times 10^{18}$ eV, with high efficiency and unprecedented statistical accuracy. The energy spectrum of Ultra-High Energy Cosmic Rays (UHECR) with inclined showers was recently measured~\\cite{FacalSanLuis:2007it}. The surface detector is an extended array of deep water Cherenkov detectors that act like volume detectors adequate for recording particles arriving at ground at all zenith angles.  As in all ground arrays, the distribution of the detector signals produced by shower particles is used to estimate shower observables such as the primary energy and to perform composition studies. As shown in previous works~\\cite{Inoue:1999cn,Ave:2000xs}, the specific characteristics of inclined showers entail that their analysis requires a different approach from the standard one for vertical showers. The study of the particle densities of the electromagnetic and muonic components at ground level becomes essential in the reconstruction~\\cite{Newton:2007qi} and analysis of events at large zenith angles.  Inclined showers have also been studied for years for other reasons. In the 70's, it was suggested that showers induced by neutrinos could be identified in the background of inclined proton or nuclei induced showers by searching for deeply penetrating inclined showers~\\cite{Berezinsky1975} which should exhibit a significant electromagnetic component at ground. In this respect the study and characterisation of the electromagnetic halo in inclined showers is also of great importance.   In this paper we have used Monte Carlo generators simulating the air shower development to obtain the particle densities induced by the electromagnetic and muonic components of inclined air showers. In contrast to other works~\\cite{Ave:1999cp,Cillis:2000xc} we have performed a comprehensive characterisation of the ratio of the electromagnetic to muonic densities as a function of distance to the shower core for different zenith angles. We have studied the dependences of this ratio on the primary energy, mass composition and hadronic interaction model assumed in the simulations. We have also studied the effect of the geomagnetic field that deviates charged particles along their paths to the detector. As a result of this study we give parameterisations of the average ratio of the EM and muonic components as a function of the distance to the shower core, shower zenith angle~($\\theta$) and azimuth angle ($\\zeta$) of the position at which the particle arrives at ground with respect to the incoming shower direction, as well as on the distance to the shower core in the shower plane~\\footnote{The shower plane is the plane transverse to the shower axis.}.  The effect of the intrinsic shower-to-shower fluctuations on the electromagnetic and muonic contributions to the density has also been studied. Our results are relevant for the reconstruction of inclined showers in ultra-high energy cosmic ray detectors such as the Pierre Auger Observatory and the Telescope Array. ",
        "conclusions": "We have studied the characteristics of the particle densities of the electromagnetic and muon components of inclined showers at the ground level using Monte Carlo simulations. We have shown that the electromagnetic component is composed of several sub-components originated in cascading and muonic processes. These different contributions to the electromagnetic component differ from each other on their behaviour with distance to the core, zenith angle and angular position ($\\zeta$).  We have studied the ratio of the electromagnetic to muon densities ($R_{\\rm EM/\\mu}$) as a function of several parameters. Firstly, we have characterised the dependence of this ratio on the distance to the core and shower zenith angle (see Fig.~\\ref{EmMuRatio}). Near the core up to $\\theta \\sim 70^{\\circ}$ its behaviour is explained by the increasing absorption of the contribution to the EM component due to $\\pi^0$ decay and beyond $70^{\\circ}$ by the dominance of the contribution to the EM component due hard muon processes. Far from the core ($> 1$ km), the ratio is compatible with an almost constant value because the electromagnetic component due to muon decay in flight dominates the electromagnetic density at ground. Then, we have studied the dependence of $R_{\\rm EM/\\mu}$ on the azimuthal position $\\zeta$. For showers with $60^{\\circ} < \\theta < 70^{\\circ}$ we have found an azimuthal asymmetry, which is mostly due to the longitudinal development effect. Moreover, we have studied the effect of the Earth magnetic field in this ratio finding that is important for showers at $\\theta \\gtrsim 84^{\\circ}$.  Considering all these dependences above and neglecting the geomagnetic field effect, we propose a parameterisation of the $R_{\\rm EM/\\mu}(r,\\theta,\\zeta)$ that predicts the simulated data to a good precision level (see Fig.~\\ref{ExampleAsyEmMuRatio}). The fits are valid for 10 EeV proton showers simulated with the hadronic model QGSJET01, and in the ranges $\\theta \\in [60^{\\circ},89^{\\circ}]$ and $r \\in [20,4000]$ m.  We have characterised the dependence of this ratio with the primary energy, the primary mass composition and the hadronic interaction model used in the simulations. The general result is that at zenith angles $ \\gtrsim 70^{\\circ}$ the ratio remains constant because only the electromagnetic halo contributes to the electromagnetic component.  Finally, we have estimated the effect of physical fluctuations on the electromagnetic and muon densities. The muon fluctuations remain roughly constant at the level of $ \\lesssim 20\\%$. While, the EM fluctuations depend on the dominance of the electromagnetic halo on the EM component, and tend to be constant $\\lesssim 20\\%$ for higher zenith angles.  In Table~\\ref{tabla1} we summarise some results for two different distances to the core and different shower zenith angles. Note that in the angular region $60^{\\circ} < \\theta < 70^{\\circ}$ the presence of the remnant of the electromagnetic shower due to cascading processes from $\\pi^0$ decay of hadronic origin entails important dependences of the ratio on the energy, mass composition and hadronic model. Therefore, this region could be the most suitable one for studies about composition and hadronic interaction models by means of the electromagnetic component in the shower at ground level. However in the region $\\theta \\gtrsim 70^{\\circ}$ the ratio is not sensitive to composition and hadronic model, and this angular region might be in principle more suitable to study the UHECR energy spectrum with inclined events detected by ground detectors.     \\begin{table}[] \\begin{center} \\begin{tabular}{||c|c|c|c|c|c|c|c||} \\hline\\hline \\small{r (m)} & \\small{$\\theta$ (deg)} & \\small{$\\Delta R_{{B}_{\\phi = 180^{\\circ}}}$} & \\small{$\\Delta R_{E=100 EeV}$} & \\small{$\\Delta R_{\\rm mass}$} & \\small{$\\Delta R_{\\rm had}$} & \\small{$\\sigma_{{\\rho_{\\rm EM}}}^{\\rm rel}$}& \\small{$\\sigma_{\\rho_{\\mu}}^{\\rm rel}$}\\\\ \\hline \\hline 100 &  60 & -2  & 83 &  -61 & 40 &  45 & 13  \\\\ \\cline{2-8} & 70 & 7  & 6 &  5 & 11 &  14 &  12  \\\\ \\cline{2-8} & 80 & 29  & -2 &  5 & 1 &  17 &  16  \\\\ \\hline 1000 &  60 & -2  & 33 &  -42 & 19 &  26 &  13  \\\\ \\cline{2-8} & 70 & 7    & 4 &  2 & 3 & 14 & 12  \\\\ \\cline{2-8} & 80 & 9  & -2 &  -2 & -1 & 17&  15  \\\\ \\hline\\hline \\end{tabular} \\caption{Summary of the dependences of the ratio of EM to muonic component on the geomagnetic field, energy, primary mass and hadronic model, using protons at $E=10^{19}$ eV simulated with the QGSJET01 model and with no geomagnetic field as reference (see Eqs.~\\ref{deltaB}, \\ref{deltaE}, \\ref{deltaM} and \\ref{deltaH} for the definitions of the different $\\Delta R$). The dependences are given at different distances from the shower axis 100 m and 1 km, and different zenith angles $60^{\\circ}$, $70^{\\circ}$ and $80^{\\circ}$. We also give the relative physical fluctuations of the EM and muonic components as obtained with Eq.~\\ref{sigmaN_shtosh}. The results are given in percentage $\\%$.\\label{tabla1}} \\end{center} \\end{table}"
    },
    "0910/0910.2935_arXiv.txt": {
        "abstract": " ",
        "introduction": "As possible hosts of the Square Kilometre Array (SKA), South Africa and Australia are building SKA Precursor arrays: MeerKAT and ASKAP, respectively. The two telescopes will complement each other well: ASKAP will have a wider field of view but a smaller frequency range and lower sensitivity, while MeerKAT will be more sensitive, have a larger frequency range, but with a smaller field of view.  MeerKAT will have additional shorter and longer baselines, giving it enhanced surface brightness sensitivity as well as astrometric capability. It is also envisaged that MeerKAT will have the capability of phasing-up array elements and will, from time to time, participate in the European, Australian and global VLBI networks.  This document gives a short overview of the expected scientific capabilities as well as the technical specifications of the MeerKAT telescope and invites the community to submit Large Project proposals that take advantage of the unique capabilities of the instrument.  In this document we give a short description of the MeerKAT array in Sect.~2. A brief summary of MeerKAT key science is given in Sect.~3. More detailed technical specifications and the array configuration are given in Sect.~4, with the proposal format and policies listed in Sect.~5. A summary is given in Sect.~6. ",
        "conclusions": "With this short discussion of the MeerKAT scientific goals and applications, we invite Large Proposals from teams prepared to help design observational programmes, and help contribute to developing simulations, and advanced observing methods and analysis techniques.  MeerKAT will be capable of very exciting science.  It will be a major pathfinder to the SKA, giving insights into many of the technical challenges of the SKA, but also giving a glimpse of the new fundamental studies that the SKA will facilitate.   \\bigskip \\emph{We acknowledge the help and stimulation given by members of the MeerKAT International Scientific Advisory Group: Bruce Bassett, Mike Garrett, Michael Kramer, Robert Laing, Scott Ransom, Steve Rawlings and Lister Staveley-Smith. We also thank Michael Bietenholz, Bradley Frank and Richard Strom for their help and contributions.}   \\newpage"
    },
    "0910/0910.1107_arXiv.txt": {
        "abstract": "We report the results of spectroscopic mapping observations carried out toward protostellar outflows in the BHR71, L1157, L1448, NGC 2071, and VLA 1623 molecular regions using the Infrared Spectrograph (IRS) of the {\\it Spitzer Space Telescope}.  These observations, covering the $5.2 - 37\\rm\\,\\mu m$ spectral region, provide detailed maps of the 8 lowest pure rotational lines of molecular hydrogen and of the [SI] 25.25$\\rm\\,\\mu m$  and [FeII] 26.0$\\rm\\,\\mu m$ fine structure lines.  The molecular hydrogen lines, believed to account for a large fraction of the radiative cooling from warm molecular gas that has been heated by a non-dissociative shock, allow the energetics of the outflows to be elucidated.  Within the regions mapped towards these 5 outflow sources, total H$_2$ luminosities ranging from 0.02 to 0.75 L$_{\\odot}$ were inferred for the sum of the 8 lowest pure rotational transitions. By contrast, the much weaker [FeII] 26.0$\\rm\\,\\mu m$ fine structure transition traces faster, dissociative shocks; here, only a small fraction of the fast shock luminosity emerges as line radiation that can be detected with {\\it Spitzer}/IRS. ",
        "introduction": "It has long been recognized that the formation of stars in molecular clouds is often accompanied by supersonic outflows (e.g. Bally \\& Lada 1983).  Where protostellar outflows strike nearby molecular material, shock waves can be created; these compress and heat the gas, and can modify its composition.  Interstellar shocks have been the subject of extensive study, both observational and theoretical (summarized, for example, in the review of Draine \\& McKee 1993).  Slow shocks -- propagating at speeds less than $\\sim 50 \\rm  \\, km\\, s^{-1}$ (in clouds with typical magnetic fields) -- leave the gas predominantly molecular in their wake, although the chemical composition can be modified by two processes: (1) the sputtering of dust grains and their icy mantles, which releases material into the gas phase; and (2) the enhancement of gas-phase chemical reactions that possess activation energy barriers; these may be negligibly slow at the typical temperatures of molecular clouds but rapid at the elevated temperatures behind shock waves.  Faster shocks can lead to the collisional dissociation of molecules and the ionization of the resultant atoms.   \\re{ The Infrared Spectrograph (IRS) of the {\\it Spitzer Space Telescope} provides a valuable tool for observing emissions from shocked interstellar gas.  With its coverage of $5.2 - 37\\rm\\,\\mu m$ spectral region, IRS has provided access to (1) the lowest 8 pure rotational transitions of molecular hydrogen (e.g.\\ Neufeld et al. 2006a, hereafter N06) -- the dominant constituent of molecular clouds; (2) pure rotational transitions of HD (e.g.\\ Neufeld et al. 2006b), H$_2$O, and OH (e.g. Melnick et al.\\ 2008); (3) rovibrational transitions of CO$_2$ and C$_2$H$_2$ (e.g. Sonnentrucker et al.\\ 2006; Sonnentrucker, Gonz\\'alez-Alfonso \\& Neufeld 2007); (4) fine structure emissions of several atoms and atomic ions (Neufeld et al.\\ 2007, hereafter N07).   Observations of shocked material at mid-infrared wavelengths offer significant advantages over ultraviolet, optical or even near-infrared studies: they are affected less by interstellar extinction, and they can probe cooler material than can be observed at shorter wavelengths.}  In this paper, we present the results of spectroscopic mapping observations carried out using the {\\it Spitzer}/IRS) toward shocked material associated with five protostellar outflows: BHR71, L1157, L1448, NGC 2071, and VLA 1623.  In \\S 2 below, we discuss the properties of the individual sources and the rationale underlying our source selection.  In \\S 3, we describe the observations, discuss our data reduction procedure, and present average spectra and spectral line maps for each source.  In \\S 4, we discuss the basic features of the H$_2$ rotational emissions and atomic fine structure emissions that we detected; here we analyse the overall energetics for each source, describe the correlations among the various detected line emissions, and compare our H$_2$ spectral line maps with continuum maps obtained previously using {\\it Spitzer}'s Infrared Array Camera (IRAC) and extracted by us from the {\\it Spitzer} archive. ",
        "conclusions": "\\subsection{H$_2$ rotational emissions}  The pure rotational transitions of H$_2$ are optically thin, and thus the level populations are readily calculable.  Our observations of the S(0) to S(7) rotational emissions provide estimates of the $J=2$ through $J=9$ column densities.  In this paper, we consider only the total H$_2$ luminosities for the entire mapped region in each object.  We defer a detailed discussion of spatial variations in the line luminosities and line ratios to a future paper, (Nisini et al.\\, \\re{ in preparation), in which we will also discuss existing spectral line maps of H$_2$ vibrational emissions.}  In Figure 16, we present rotational diagrams based upon the H$_2$ luminosities given in Table 3, corrected for extinction using the extinction estimates given in Table 1.  Here, we plot the logarithm of the total number of molecules in each state, divided by the degeneracy of that state, as a function of the energy.  Diamonds indicate the observed values, and the solid lines are results obtained from the fitting procedure described below.   While material in thermal equilibrium at a single temperature would exhibit a linear behavior in such a plot, the observed rotational diagrams -- like those obtained for outflow sources by N06 and by Maret et al.\\ (2009), and for supernova remnants by N07 -- depart from that behavior in two respects.  First, there is a positive curvature to the rotational diagram, indicative of the presence of multiple temperature components.  Second, except toward L1448, there is a characteristic zigzag behavior, most notably for the states of lowest energy, in which the states of even rotational quantum number $J$ lie systematically above those of odd $J$.  As in N06, N07, and Maret al.\\ (2009), we attribute this second behavior to an ortho-to-para ratio -- smaller than the equilibrium value of 3 -- that is the relic of an earlier period when the gas was considerably cooler than it is now.  We have modeled the H$_2$ rotational diagrams in Figure 16, using a method similar to that described by Neufeld and Yuan (2008; hereafter NY08).   As in NY08 \\re{ (their eqn.\\ 3)}, we assume a power law distribution of gas temperatures with the number of molecules at temperature between $T$ and $T+dT$ assumed proportional to $T^{-b}$.  We assume a temperature range extending from $T_{min} = 100$~K \\footnote{Recognizing that gas in the 100 -- 300~K can contribute significantly to the pure rotational line emission that we detect with {\\it Spitzer}, we have reduced $T_{min}$ from 300~K in NY08 to 100~K in the present work.} to $T_{max}=4000$~K.  Solving the equations of statistical equilibrium for the H$_2$ level populations, and adopting the molecular data cited in NY08, we obtained predictions for the pure rotational line luminosities as a function of the power law index, $b$, the total mass of hydrogen molecules warmer than $T_{min}$, $M_{\\rm H2}(\\ge T_{min})$, and the H$_2$ density, $n({\\rm H}_2)$. \\re{ The excitation of pure rotational emissions is dominated by collisions with H$_2$ (see NY08), in contrast to the case of rovibrational emissions, where excitation by atomic hydrogen presents an additional complication that we will discuss in a future paper (Nisini et al.\\ 2009, in preparation)}.  We have extended the treatment of NY08 by considering the H$_2$ ortho-to-para ratio along with $b$, $M_{\\rm H2}(\\ge T_{min})$ and $n({\\rm H}_2)$.  Here, we introduce two additional parameters: the initial ortho-to-para ratio, OPR$_0$, and the time period, $\\tau$, for which the gas temperature has been elevated.  If reactive collisions with atomic hydrogen are the dominant mechanism for the interconversion of ortho- and para-H$_2$, then the ortho-to-para ratio after time $\\tau$ is given by N06\\footnote{Note that the expression given by N06 erroneously omits $k_{\\rm OP}$ from the term $[k_{\\rm PO}+k_{\\rm OP}]$, resulting in a 25$\\%$ error.}: $${\\rm OPR(\\tau) \\over 1+OPR(\\tau)} =  {\\rm OPR_0 \\over 1+OPR_0} \\, e^{-n({\\rm H})\\, \\tau \\, [k_{\\rm PO}+k_{\\rm OP}]} +  {\\rm OPR_{LTE} \\over 1+OPR_{LTE}} \\, \\biggl( 1- e^{-n(H)\\, \\tau \\, [k_{\\rm PO}+k_{\\rm OP}]}\\, \\biggr), \\eqno(1)$$ where $n({\\rm H})$ is the atomic hydrogen density, $k_{\\rm PO}$ is the rate coefficient for para-to-ortho conversion, estimated as $8 \\times 10^{-11} \\exp(-3900/T) \\rm \\, cm^3 \\, s^{-1}$ by Schofield et al.\\ (1967), based upon the laboratory experiments of (Schulz \\& Le Roy 1965), and $k_{\\rm OP} \\sim k_{\\rm PO}/3$ is the rate coefficient for ortho-to-para conversion.  Thanks to the strong dependence of $k_{PO}$ upon the gas temperature, the ortho-to-para ratio is similarly temperature-dependent, the warmer gas components being closer to equilbrium than the cooler ones. This effect is evident in Figure 16, the degree of zigzag being least for the highest energy states that are excited primarily in the warmest gas.  We have adjusted the five parameters $b$, $M({\\rm H}_2)$, $n({\\rm H}_2)$, OPR$_0$, and $n({\\rm H})\\tau$ in order to obtain the best fit to the data.  Here, we minimize the $\\chi^2$ for the S(1) through S(7) rotational lines, assuming the same fractional uncertainty for each data point.  We exclude the H$_2$ S(0) line flux from the analysis, because its value is rather uncertain, particularly in regions of strong continuum intensity where the line-to-continuum ratio becomes small. The best fit parameters are given in Table 3.  The inferred masses of warm\\footnote{ \\re{ In this paper, we use the word ``warm\" to describe H$_2$ at temperatures $\\ge 100$~K.  This, of course, includes gas at lower temperatures than that which can be probed by means of traditional, ground-based studies of vibrationally-excited H$_2$.}}  H$_2$ (i.e.\\ at $T \\ge 100$~K) range from $1.3 \\times 10^{-3} M_\\odot$ in VLA 1623 to $0.18 M_\\odot$ in NGC 2071.  These values exclude the mass contribution from helium and apply only to the regions mapped by IRS.  The best-fit power-law indices, $b$, range from 2.2 to 3.3.   These values lie somewhat below the typical values determined by NY08 in the SNR IC443C, and somewhat below the value $b = 3.8$ expected for a paraboloidal bow shock that is dissociative at its apex.  The best fit H$_2$ densities range from $3000$ to $7000\\,\\rm cm^{-3}$.  It should be noted, however, that the range of acceptable values for $b$ and $n({\\rm H}_2)$ can be large, and that these two parameters are somewhat degenerate in the following sense: increasing $b$ tends to decrease the predicted relative strength for the higher $J$ transitions, whilst decreasing $n(\\rm H_2)$ has a similar effect (because the level populations become increasingly subthermal).  Thus, an increase in one of these assumed parameters can partially compensate for a decrease in the other.  The correlated uncertainties in $b$ and $n(\\rm H_2)$ are presented graphically in Figure 17, which shows contours of $\\chi^2$ in the $b$--$n(\\rm H_2)$ plane.  Crosses indicate the minimum $\\chi^2$ (i.e. the best fit), while the two contours represent the 68$\\%$ and 95$\\%$ confidence limits on each parameter.  With the exception of L1448, for which the data are consistent with an OPR of 3, the objects require an initial OPR less than 3.  In two objects, L1157 and NGC 2071, the current OPR is clearly larger for the higher-$J$ states than for the lower-$J$ states.  In the context of our model for the evolution of the OPR (eqn.\\ 1), both cases imply a lower limit $\\sim 10^3 \\,[\\rm{cm}^{-3}/n({\\rm H})]$~yr on the time period, $\\tau$, for which the gas has been warm.  \\subsection{Fine structure emissions}  The Fe$^+$ $\\,\\,\\rm ^6D_{7/2} - ^6D_{9/2}  \\,\\,\\,\\,25.988 \\,\\mu$m and S$\\,\\,\\,\\,\\rm ^3P_{1} -^3P_{2} \\,\\,\\,\\, 25.249 \\,\\mu$m fine structure lines are readily detectable in the LH spectra of all five outflows sources, and are strong enough to map.  However, these transitions make a negligible contribution to the total 5.2 -- 37 $\\mu$m line luminosity.  The [SI] line luminosity ranges from $\\sim$ 1 to 2 $\\%$ of the total H$_2$ S(0) -- S(7) luminosity.  For [FeII], the corresponding range is wider ~ $\\sim$ 0.15 to 2$\\%$.  In the case of [FeII], these percentages are much smaller than those attained in some of the supernova remnants (SNR) that we observed previously (N07).  In the SNR 3C391, for example, the [FeII] $25.988 \\,\\mu$m line flux exceeded 10$\\%$ of the total H$_2$ S(0) -- S(7) flux.  Furthermore, fine structure emissions from more highly ionized species -- Ne$^+$, Ne$^{++}$, and S$^{++}$ for example -- are readily detectable in SNR but not in protostellar outflow sources.  These differences suggest that higher velocity shocks are more prevalent in supernova remnants than around outflow sources.  A visual inspection of Figures 11 -- 15 suggests that the [SI] emissions are typically more closely correlated with the H$_2$ emissions than are the [FeII] emissions.  This impression is borne out by a correlation analysis.  Comparing the H$_2$~S(5) and [SI] intensities, spatial pixel by spatial pixel, we obtained linear correlation coefficients of 0.81, 0.78, 0.42, 0.80, and 0.85 respectively for BHR71, L1157, L1448, NGC 2071 and VLA 1623.  Comparing the H$_2$~S(5) and [FeII] intensities, the corresponding results are invariably smaller: 0.55, 0.48, 0.23, 0.72, and 0.78.  This behavior is consistent with our previous study of supernova remnants, in which a principal component analysis suggested that the [SI] emitting region was largely cospatial with the H$_2$ emitting region, while emissions from ionized species such as Fe$^+$ originated in a separate gas component.  This segregation is presumably a consequence of the different shock velocities needed to generate [SI], H$_2$, and [FeII] emissions.  The first two arise in slow non-dissociative shocks that create very little ionization, while the third arises in faster shocks that are dissociative and ionizing.  According to theoretical models (e.g.\\ Allen et al.\\ 2008), the emission from dissociative shocks is dominated by optical and ultraviolet line emissions, with infrared fine structure emissions accounting for only a small fraction of the overall shock luminosity.  \\subsection{Source energetics and comparison with IRAC maps}  The wavelength coverage of {\\it Spitzer} is well suited to observing the radiative cooling of non-dissociative molecular shocks.  The study of Kaufman and Neufeld (1996) indicates that over a wide range of shock parameters, pure rotational emissions from H$_2$ are key coolants of non-dissociative shocks.  For shocks of velocity in the range 20 to 40~km~s$^{-1}$ and with a preshock H$_2$ density of $10^4$~cm$^{-3}$, for example, between 45 and 69$\\%$ of the shock luminosity is predicted to emerge in pure rotational emission from H$_2$.  Moreover, almost all of the H$_2$ pure rotational emission emerges in the S(0) - S(7) lines accessible to {\\it Spitzer}.    The total H$_2$ S(0) -- S(7) luminosities listed in Table 1 range from 0.02 to 0.75 $\\, L_\\odot$.   Taking these luminosities as lower limits for the total mechanical luminosity dissipated in the shocks, adopting 20~km~s$^{-1}$ as a typical shock velocity, $v_s$, and setting the mechanical luminosity equal to  $\\onehalf \\dot{M} v_s^2$, we obtain lower limits on the mass flow rates through the shock, $\\dot{M}$,  ranging from $\\sim 10^{-6}$ to $\\sim 2 \\times 10^{-5} M_\\odot \\rm \\, yr^{-1}$.  In some objects, H$_2$ line emissions can contribute significantly to the intensities measured using broadband mid-IR photometry (e.g. Reach et al.\\ 2006, NY08).  For the supernova remnant IC443, NY08 found that the mid-IR spectral region covered by IRAC was dominated by H$_2$ line emissions.  In particular, the intensities measured in the 5.8~$\\mu$m and 8~$\\mu$m IRAC channels (Bands 3 and 4), were accounted for almost entirely by the IRS-measured  H$_2$ S(4) - S(7) line intensities.  Given an excitation model for H$_2$, the 3.6 $\\mu$m and 4.5 $\\mu$m IRAC intensities could similarly be explained as resulting from H$_2$ emissions, primarily the S(9) pure rotational emission for the 4.5 $\\mu$m band and the v = 1 -- 0 Q(5) rovibrational line for the 3.6 $\\mu$m band.  For the protostellar outflow sources studied here, we can evaluate the contribution of the H$_2$ S(4) -- S(7) emissions to the broadband intensities measured with IRAC.  In Table 3, we list the fraction of the intensities measured in the 5.8~$\\mu$m and 8~$\\mu$m IRAC bands that is attributable to H$_2$ S(4) and S(5) (in the case of the 8~$\\mu$m band) or H$_2$ S(6) and S(7) (in the case of the 5.8~$\\mu$m band). The contribution is clearly largest in BHR71 and L1157, for which the computed contribution to the 8~$\\mu$m band actually exceeds 100$\\%$, a discrepancy that presumably reflects flux calibration errors that are within expectations.  Even in these objects, however, the 5.8~$\\mu$m band intensity is 2--3 times stronger than what can be attributed to H$_2$, and in the three other objects, the contribution of H$_2$ to either the 5.8~$\\mu$m or 8.0~$\\mu$m band is small, at least averaged over the entire mapped region.  Nevertheless, the IRAC maps shown in Figures 11 - 15 can exhibit clear morphological similarities to the H$_2$ line maps, even in objects such as L1448 where the overall contribution of H$_2$ is small.  This reflects the fact that regions of diffuse emission can be H$_2$-dominated even if compact continuum sources account for most of the IRAC intensity when averaged over the entire region that we have mapped.  Finally, we note that in the case of VLA 1623 and NGC 2071, the contributions listed in Table 3 reflect an important selection effect: we specifically avoided observing H$_2$ emissions in those regions known to exhibit the strongest continuum emission, because they would have saturated the IRS detectors.  Thus, had we been able to map the entire object with IRS, the fractional contributions of H$_2$ listed in Table 3 would have been even smaller, probably by a large factor.  In summary, a comparison of the IRS and IRAC data indicates that the 8$\\mu$m IRAC channel can be dominated by H$_2$ emissions for a protostellar outflow, as it was for the supernova remnant IC443 discussed by NY08, but that the H$_2$ contribution can also be negligible, as in NGC 2071 and VLA 1623."
    },
    "0910/0910.1041_arXiv.txt": {
        "abstract": "We consider the renormalization group improvement in the theory of the Standard Model (SM) Higgs boson playing the role of an inflaton with a strong non-minimal coupling to gravity. It suggests the range of the Higgs mass $135.6\\; {\\rm GeV} \\lesssim M_H\\lesssim 184.5\\;{\\rm GeV}$ entirely determined by the lower WMAP bound on the CMB spectral index. We find the phenomenon of asymptotic freedom induced by this non-minimal curvature coupling, which brings the theory to the weak coupling domain. Asymptotic freedom fails at the boundaries of this domain, which makes the SM phenomenology sensitive to the current cosmological data and thus suggests future more precise CMB measurements as a SM test complementary to the LHC program. By using a concept of field-dependent cutoff we also show naturalness of the gradient and curvature expansion in this model within a conventional perturbation theory range of SM. ",
        "introduction": "The hope for the forthcoming discovery of the Higgs boson at LHC revives the attempts of constructing a particle model accounting also for an inflationary scenario and its observable CMB spectrum. An obvious rationale behind this is the anticipation that cosmological observations can comprise Standard Model (SM) tests complimentary to collider experiments. While the relatively old work \\cite{we-scale} had suggested that due to quantum effects inflation depends not only on the inflaton-graviton sector of the system but rather is strongly effected by the GUT contents of the particle model, the series of papers \\cite{BezShap,we,BezShap1,Wil,BezShap3} transcended this idea to the SM ground with the Higgs field playing the role of an inflaton. This has recovered interest in a once rather popular \\cite{non-min,we-scale,BK,KomatsuFutamase} model with the Lagrangian of the graviton-inflaton sector \\begin{eqnarray} &&{\\mbox{\\boldmath $L$}}(g_{\\mu\\nu},\\varPhi)= \\frac12\\left(M_P^2+\\xi|\\varPhi|^2\\right)R -\\frac{1}{2}|\\nabla\\varPhi|^{2} -V(|\\Phi|),                     \\label{inf-grav}\\\\ &&V(|\\Phi|)= \\frac{\\lambda}{4}(|\\varPhi|^2-v^2)^2,\\,\\,\\,\\, |\\varPhi|^2=\\varPhi^\\dag\\varPhi, \\end{eqnarray} where $\\varPhi$ is a scalar field whose expectation value plays the role of an inflaton and which has a strong non-minimal curvature coupling with $\\xi\\gg 1$. Here, $M_P=m_P/\\sqrt{8\\pi}\\approx 2.4\\times 10^{18}$ GeV is a reduced Planck mass, $\\lambda$ is a quartic self-coupling of $\\varPhi$, and $v$ is a symmetry breaking scale.  The motivation for this model was  based on the observation \\cite{non-min} that the problem of an exceedingly small quartic coupling $\\lambda\\sim 10^{-13}$, dictated by the amplitude of primordial CMB perturbations \\cite{CMB}, can be solved by using a non-minimally coupled inflaton with a large value of $\\xi$. Later this model with the GUT-type sector of matter fields was used to generate initial conditions for inflation \\cite{we-scale} within the concept of the no-boundary \\cite{noboundary} and tunneling cosmological state \\cite{tunnel}. The quantum evolution with these initial data was considered in \\cite{BK,efeqmy}. There it was shown that quantum effects are critically important for this scenario.  A similar theory but with the SM Higgs boson $\\varPhi$ playing the role of an inflaton instead of the abstract GUT setup of \\cite{we-scale,BK} was suggested in \\cite{BezShap}. This work advocated the consistency of the corresponding CMB data with WMAP observations at the tree-level approximation of the theory, which was extended in \\cite{we} to the one-loop level. This has led to the lower bound on the Higgs mass $M_H\\gtrsim 230$ GeV, originating from the observational restrictions on the CMB spectral index \\cite{we}. However, this conclusion which contradicts the widely accepted range 115 GeV$\\leq M_H\\leq 180$ GeV did not take into account $O(1)$ contributions due to renormalization group (RG) running, which qualitatively changes the situation. This was nearly simultaneously observed in \\cite{GB08,BezShap1,Wil} where the RG improvement of the one-loop results of \\cite{we} predicted the Higgs mass range almost coinciding with the conventional one.  Here we suggest the RG improvement of our one-loop results in \\cite{we} and find the CMB compatible range of the Higgs mass, both boundaries of this range being determined by the lower WMAP bound on the CMB spectral index, $n_s\\gtrsim 0.94$ rather than by perturbation theory arguments. This makes the phenomenology of this gravitating SM essentially more sensitive to the cosmological bounds than in \\cite{BezShap1,Wil,BezShap3} and makes it testable by the CMB observations. The mechanism underlying these conclusions and explaining the efficiency of the perturbation theory at the inflationary scale is a phenomenon analogous to asymptotic freedom. Below we briefly present these results, whereas their detailed derivation and the discussion of their limitations can be found in \\cite{RGHiggs}. ",
        "conclusions": "We have found that the considered model looks remarkably consistent with CMB observations in the Higgs mass range \\begin{equation} 135.6\\; {\\rm GeV} \\lesssim M_H \\lesssim 184.5\\;{\\rm GeV},                \\label{range} \\end{equation} which is very close to the widely accepted range dictated by electroweak vacuum stability and perturbation theory bounds.  This result is based on the RG improvement of our analytical results in \\cite{we}. We have completely recovered the analytic formalism of \\cite{we} for all inflation parameters, which only gets modified by the RG mapping between the coupling constants at the EW scale and those at the scale of inflation. A peculiarity of this formalism is that for large $\\xi\\gg 1$ the effect of SM phenomenology on inflation is universally encoded in one quantity -- the anomalous scaling ${\\mbox{\\boldmath $A_I$}}$. It was earlier suggested in \\cite{we-scale} for a generic gauge theory, and in SM it is dominated by contributions of heavy particles -- ($W^\\pm$, $Z$)-bosons, top quark and Goldstone modes. This quantity is forced to run in view of RG resummation of leading logarithms, and this running raises a large negative EW value of ${\\mbox{\\boldmath $A_I$}}$ to a {\\em small negative} value at the inflation scale. Ultimately this leads to the admissible range of Higgs masses  very close to the conventional SM range.  This mechanism can be interpreted as asymptotic freedom, because ${\\mbox{\\boldmath $A_I$}}/64\\pi^2$ determines the strength of quantum corrections in inflationary dynamics \\cite{BK,we}. Usually, asymptotic freedom is associated with the asymptotic decrease of some running coupling constants to zero. Here this phenomenon is trickier because it occurs in the interior of the range  and fails near its lower and upper boundaries. Quantum effects are small only in the middle part of (\\ref{range}) with a moderately small $\\lambda$ where $n_s$ is close to the ``classical'' limit $1-2/N\\simeq 0.967$ for $x\\equiv N{\\mbox{\\boldmath $A$}}/48\\pi^2\\ll 1$. Thus, the original claim of \\cite{BezShap1} on smallness of quantum corrections is right, but this smallness, wherever it takes place, is achieved via a RG summation of big leading logarithms.  Qualitatively our main conclusions are close to those of \\cite{Wil} and \\cite{BezShap3}, though our bounds on the SM Higgs mass are much more sensitive to the CMB data. Therefore, the latter can be considered as a test of the SM theory complimentary to LHC and other collider experiments. The source of this difference from \\cite{Wil,BezShap3} can be ascribed to the gauge and parametrization (conformal frame) dependence of the off-shell effective action (along with the omission of Goldstone modes contribution in \\cite{Wil}) -- an issue which is discussed in much detail in \\cite{RGHiggs} and which is expected to be resolved in future publications.  We have also shown the naturalness of the gradient and curvature expansion in this model, which is guaranteed within the conventional perturbation theory range of SM, $\\lambda/16\\pi^2\\ll 1$, and holds in the whole range of the CMB compatible Higgs mass (\\ref{range}) -- the latter property being a consequence of the asymptotic freedom of the above type. This result is achieved by the background field resummation of weak field perturbation theory leading to the replacement of the fundamental Planck mass in the known cutoff $4\\pi M_P/\\xi$ \\cite{BurgLeeTrott,BarbEsp} by the effective one. Partly (modulo corrections to inflaton potential, which are unlikely to spoil its shape) this refutes objections of \\cite{BurgLeeTrott,BarbEsp} based on the analysis of scattering amplitudes in EW vacuum background. Smallness of the cutoff in this background does not contradict physical bounds on the Higgs mass originating from CMB data for the following reasons. Determination of $M_H$ of course takes place at the TeV scale much below the non-minimal Higgs cutoff $4\\pi M_P/\\xi$, whereas inflationary dynamics and CMB formation occur for $\\lambda/16\\pi^2\\ll 1$ below the {\\em running} cutoff $\\Lambda(\\varphi)=4\\pi\\varphi/\\sqrt\\xi$. It is the phenomenon of inflation which due to exponentially large stretching brings these two scales in touch and allows us to probe the physics of underlying SM by CMB observations at the 500 Mpc wavelength scale.  To summarize, the inflation scenario driven by the SM Higgs boson with a strong non-minimal coupling to curvature looks very promising. This model supports the hypothesis that the appropriately extended Standard Model can be a consistent quantum field theory all the way up to quantum gravity and perhaps explain the fundamentals of all major phenomena in early and late cosmology \\cite{nuMSM,dark}. Recently, an interesting question about the possibility of a supersymmetric generalizations of the model was considered \\cite{supersym}. In principle, the study of this question can shed a new light on the problem of the range  of the validity of the model. Ultimately, it will be the strongly anticipated discovery of the Higgs particle at LHC and the more precise determination of the primordial spectral index $n_s$ by the Planck satellite that might decide the fate of this model."
    },
    "0910/0910.3218.txt": {
        "abstract": "By cross matching blue objects from SDSS with GALEX and the astrometric catalogues USNO-B1.0, GSC2.3 and CMC14, 64 new dwarf nova candidates with one or more observed outbursts have been identified. 14 of these systems are confirmed as cataclysmic variables through existing and follow-up spectroscopy. A study of the amplitude distribution and an estimate of the outburst frequency of these new dwarf novae and those discovered by the Catalina Real-time Transient Survey (CRTS) indicates that besides systems that are faint because they are farther away, there also exists a population of intrinsically faint dwarf novae with rare outbursts. ",
        "introduction": "One of the astrophysically interesting evolutionary phases a close binary star can go through, is that of a cataclysmic variable (CV), in which a main-sequence star filling its Roche lobe loses mass to a white dwarf.  In the case the white dwarf has no or only a weak magnetic field, an accretion disc forms around it.  Because of thermal instabilities in the accretion disc some of these systems, the dwarf novae, occasionally outburst and become several magnitudes brighter for a few days to a few weeks at most \\citep{meyer+meyer-hofmeister84-1, osaki96-1, lasota01-1}. Theoretical studies of the orbital period distribution of these dwarf novae predicted a large number of systems with an orbital period near the period minimum \\citep{kolb93-1, howelletal01-1}. Until recently the predicted accumulation of CVs near the period minimum was not what was observed \\citep{kolb+baraffe99-1, willemsetal05-1}. Furthermore there is a discrepancy between the theoretical period minimum of $\\approx$ 70 minutes and the observed one between 80 and 86 minutes \\citep{kolb+baraffe99-1}.  The Sloan Digital Sky Survey (SDSS) has led to the spectroscopic identification of a large sample of faint CVs \\citep[see][and earlier papers in the series]{szkodyetal09-1}. Follow-up studies reveal that this faint SDSS CV population differs in many aspects from the previously known brighter systems, and that its period distribution is in closer agreement with the theoretical predictions \\citep{gaensickeetal09-2}, although the discrepancy between the observed and the theoretically predicted orbital period minimum remains.  Also the observed ratio of short to long period CVs still disagrees with the model predictions \\citep{aungwerojwit06-1,pretoriusetal07-1,pretorius+knigge08-1}. While SDSS has made already a substantial contribution to observational population studies of CVs, the completeness of the currently known CV sample remains rather uncertain.  Prior to SDSS, a large fraction of all CVs in general, and the majority of dwarf novae in particular, have been discovered serendipitously by their large amplitude variability, favouring those systems with frequent outbursts \\citep{gaensicke05-1}. The majority of dwarf novae are expected to have evolved to short periods near the period minimum, and to have a low mass transfer rate. Consequently, they have infrequent outbursts, and the probability for a serendipitous discovery is lower than for systems with longer orbital periods and higher mass transfer rates. Historically, a large fraction of these serendipitous discoveries were made by amateur observers, which implies that the current sample of dwarf novae is biased towards relatively bright objects. In addition, it is very difficult to assess the spatial and temporal coverage of their searches for transient objects. The Catalina Real-time Transient Survey\\footnote{http://voeventnet.caltech.edu/feeds/Catalina.shtml} \\citep[CRTS;][] {drakeetal09-1}, based on data from the Catalina Sky Survey (CSS), is now routinely discovering dwarf novae.  Because this is a much more systematic survey, which also goes to fainter magnitudes, it is expected that it will eventually provide strong constraints on the total number and the magnitude distribution of dwarf novae.  Here, we cross-match a number of large surveys to find faint outbursting dwarf novae, and make use of CRTS light curves to compare the properties of the previously known dwarf novae, those identified spectroscopically by SDSS, and the ones discovered in this paper.    %__________________________________________________________________ ",
        "conclusions": "In this paper 64 new CV candidates, mostly dwarf novae, are presented, resulting from data mining between existing photometric and astrometric catalogues. 16 of these systems are unmistakably confirmed as CVs through follow-up spectroscopy.  It is shown that a large number of faint CVs are missed by surveys like SDSS, but can be picked up by either transient surveys like CRTS, or by archival plate studies.  The importance of having historical plates available electronically and measured, cannot be overestimated in this regard. Even including the recent CRTS discoveries, it is estimated that fewer than half the dwarf novae with a quiescent $r \\le 20$ are currently known in the SDSS fields, and substantially fewer of the fainter ones.  It appears that CVs will constitute a significant source of transients for future surveys such as LSST and Pan-STARRS.  Studying amplitude and outburst frequency suggests that the new faint dwarf novae consist of systems both similar to the previously known ones but simply farther away, and of intrinsically fainter systems. The data that will become available in the near future from CRTS and the planned surveys, will give an opportunity to study the latter population in much more detail.  Both SDSS\\,J093946.03+065209.4 and SDSS\\,J162900.55 +341022.0, two of the dwarf novae candidates found, have unusual spectra for cataclysmic variables and deserve to be studied in more detail.  % %______________________________________________________________"
    },
    "0910/0910.2986_arXiv.txt": {
        "abstract": "We report the detection of the \\paa\\ emission line in the $z=2.515$ galaxy \\target\\ using \\spitzer\\ spectroscopy.  \\target\\ is a sub-millimeter--selected infrared (IR)--luminous galaxy maintaining a high star--formation rate (SFR), with no evidence of an AGN from optical or infrared spectroscopy, nor X-ray emission.   This galaxy is lensed gravitationally by the cluster Abell 2218, making it accessible to \\spitzer\\ spectroscopy.  We measure a line luminosity, $L(\\paa)=(2.05\\pm0.33)\\times 10^{42}$~erg~s$^{-1}$, corrected for gravitational lensing.  Comparing the \\ha\\ and \\paa\\ luminosities, we derive a nebular extinction, $A(V)=3.6\\pm0.4$~mag.  The dust--corrected luminosity, $L(\\paa)=(2.57\\pm0.43)\\times10^{42}$~erg~s$^{-1}$, corresponds to an ionization rate, $Q_0=(1.6\\pm0.3)\\times10^{55}$ $\\gamma$ s$^{-1}$. The instantaneous SFR is $\\psi=171\\pm28$~\\msol~yr$^{-1}$, assuming a Salpeter--like initial mass function from 0.1 to 100~\\msol\\ yr$^{-1}$. The total IR luminosity derived using 70, 450, and 850~\\micron\\ data is $\\lir = (5-10)\\times 10^{11}$~\\lsol, corrected for gravitational lensing.  This corresponds to $\\psi=90-180$~\\msol\\ yr$^{-1}$, where the upper range is consistent with that derived from the \\paa\\ luminosity.  While the $L(8\\micron)/L(\\paa)$ ratio is consistent with the extrapolated relation observed in local galaxies and star--forming regions,  the rest--frame 24~\\micron\\ luminosity is significantly lower with respect to local galaxies of comparable \\paa\\ luminosity. Thus,  \\target\\ arguably lacks a warmer dust component ($T_D\\sim70$~K), which is associated with deeply embedded star formation, and which contrasts with local galaxies with comparable SFRs.  Rather, the starburst in \\target\\  is consistent with star--forming local galaxies with intrinsic luminosities, $\\lir\\approx10^{10}$~\\lsol, but ``scaled--up'' by a factor of $\\sim$10--100. ",
        "introduction": "There has been a growing industry of deep, multiwavelength surveys, devoted to studying star formation and evolution of galaxies. Studies of data from these surveys have concluded that the global star--formation rate (SFR) density reached a maximum between $z$$\\sim$1.5--3 \\citep[\\eg,][and references therein]{hop04}, during the period of rapid growth in the stellar mass density \\citep[\\eg,][]{dic03,rud03,rud06}.  However, the study of SFRs in these multiwavelength datasets requires estimates of the SFR using available observables.  Common SFR indicators for high--redshift galaxies  are  the rest--frame UV luminosity, especially at $z>1$ when the UV shifts into the optical bands \\citep[\\eg,][]{mad96,ste99,gia04,bou06,saw06,red08},  the far--IR luminosity inferred from sub--mm emission \\citep[\\eg,][]{blain02}, the mid--IR luminosity from \\spitzer/24~\\micron\\ \\citep[\\eg,][]{per05,cap06,pap06,red06,webb06,dad07,franx08}, and the $\\ha$--emission line \\citep[\\eg,][]{erb06,bou07,kriek08}.     Each of these SFR indicators has inherent uncertainties \\citep[\\eg,][and references therein]{ken98,hop06}, and some that are unique to observations at these redshifts (e.g., uncertainties in the luminosity--temperature dependence on the sub-mm--\\lir\\ conversion, Chapman et al.\\ 2005; Pope et al.\\ 2006; template incompleteness and photometric redshift uncertainties on the 24~\\micron--\\lir\\ conversion, Papovich et al.\\ 2006, Reddy et al.\\ 2006). Nevertheless, several studies have compared SFRs derived from the UV, 24~\\micron, far--IR, and sub--mm (and less--understood SFR tracers, such as radio and X--ray emission), concluding broadly that while they are consistent statistically, there is large scatter \\citep[\\eg,][]{dad05,dad07,red06,pope06,pap07}.  There have been few comparisons of direct, robust SFR indicators in individual galaxies at high redshifts (e.g., Siana et al.\\ 2008; references above).   SFR indicators are normally calibrated against \\ion{H}{1} recombination lines \\citep[\\eg,][]{ken98}.  The primary advantage of \\ion{H}{1} recombination lines is that they effectively re-emit the stellar luminosity from hydrogen--ionizing photons.  They are strong in star--forming regions, and they trace directly emission by massive stars.  The physics of \\ion{H}{1}  emission in star--forming regions is relatively well understood (Osterbrock 1989), and they exhibit only a weak dependence on electron density and temperature.  Of the \\ion{H}{1} lines, \\paa\\ at $\\lambda=1.8751$~\\micron\\ is an ideal practical SFR indicator.   This line is exceptionally strong in star--forming regions (1~\\msol\\ yr$^{-1}$ corresponds to $L[\\paa] = 1.48 \\times 10^{40}$ erg s$^{-1}$; Alonso--Herrero et al.\\ 2006). Furthermore, while recombination lines are subject to strong attenuation, \\paa\\ suffers minimal extinction due to its longer wavelength.  Recent studies of star formation in nearby galaxies and star--forming regions within galaxies rely on the \\paa\\ luminosity to interpret and to calibrate other SFR indicators, including the infrared emission \\citep[\\eg,][]{cal05,cal07,alo06,ken07,dia08}.  For extremely  dusty, star--forming IR luminous galaxies  with $A(V) \\sim$10--30~mag (e.g., Murphy et al.\\ 2001; Dannerbauer et al.\\ 2005; Armus et al.\\ 2007), the extinction at \\paa\\ is $\\lsim $2~mag, and optically thin.  Thus, even in the most obscured, luminous star--forming regions, the \\paa\\ line traces the intrinsic ionization rate, and thus the SFR.  While a few studies have compared SFR indicators in distant galaxies to the \\ha\\ emission line (see references above),  no attempt to use \\paa\\ as a SFR indicator has yet been attempted in any galaxy of significant redshift.  At $z$$\\gsim$2 the \\paa\\ $\\lambda 1.875$~\\micron\\ line shifts to $\\lambda_\\mathrm{obs} > 5.2$~\\micron, and is accessible to the Infrared Spectrograph (IRS) on--board \\spitzer, which covers the wavelength range of 5.2--38~\\micron. However, star--forming galaxies at $z$$>$2 are too faint intrinsically for observations with IRS at the expected wavelength for \\paa\\ (5.4--8~\\micron), except for rare, extremely luminous objects, in which AGN likely dominate the \\ion{H}{1} line emission (e.g., Brand et al.\\ 2006), complicating any SFR--tracer comparisons.  We have initiated a program using \\spitzer/IRS to study the \\paa\\ emission in 12 galaxies at $z > 2$ gravitationally lensed by foreground galaxies or clusters of galaxies.  Here, we discuss the detection of \\paa\\ in a star--forming galaxy at $z=2.515$, gravitationally lensed by the rich cluster Abell~2218 \\citep{kne04}.  This is the highest--redshift object with a measurement of the \\paa\\ emission line yet obtained, and thus we are able to study SFR indicators in a more ``typical'' star-forming galaxy directly at high redshift.  The outline for the rest of this paper is as follows.  In \\S~2, we summarize the properties of \\target.  In \\S~3, we discuss the \\spitzer\\ observations and data reduction.   In \\S~4, we discuss the analysis of the spectroscopic and imaging data to derive rest--frame luminosities.  In \\S~5, we compare the rest--frame luminosities of \\target\\ to local samples, and discuss the utility of the various quantities as SFR indicators.  In \\S~6, we present our conclusions. In the Appendix, we use stellar population models to analyze the rest--frame UV to near-IR spectral energy distribution (SED) of \\target, and make some conclusions on the nature of the stellar populations in this galaxy.      To derive physical quantities we use a cosmological model with $H_0=70$~km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{M,0} = 0.3$, $\\Omega_{\\Lambda,0}=0.7$. ",
        "conclusions": "\\label{section:discussion}  \\begin{figure*}[t] \\epsscale{1.15} \\plottwo{f7a.eps}{f7b.eps} \\epsscale{1.0} \\caption{ \\paa\\ luminosity versus the monochromatic luminosity at rest--frame 24~\\micron.  The left panel shows $L(\\paa)$ versus $L(24\\micron)$ and the right panel shows $L(\\paa)$ versus the ratio $L(24\\micron)/L(\\paa)$.  In both panels the symbols and lines are the same.   The measured value for \\target\\ is shown as the large, red--filled circle.  The light-red--shaded area indicates the error bar on the ratio for \\target.  Note that $L(24\\micron)$ for \\target\\ may be lower by $\\approx$40\\% owing to crowded photometry (see \\S~3.2).    Squares show local ultraluminous IR galaxies (ULIRGs, $\\lir \\ge 10^{12}$~\\lsol) from \\citet[D05]{dan05}, with $L(24\\micron)$ estimated using the method discussed in the text.  Diamonds show local luminous IR galaxies (LIRGs, $\\lir = 10^{11} - 10^{12}$~\\lsol) from the sample of \\citet[AH06]{alo06}; the diamond indicated by the circle corresponds to the galaxy IC~860, which is problematic according to A06.  The range of the ordinate in the right panel excludes IC~860.  Triangles show local star--forming H~II regions in M51 \\citet[C05]{cal05}. Stars show starburst and low--metallicity galaxies from \\citet[C07]{cal07} and \\citet[E05]{eng05}.     The short--dashed line shows the best fit relationship to local luminous IR galaxies from \\citet{alo06}.  The long--dashed line shows the best-fit to local galaxies and star--forming regions from \\citet{cal07}.   The short--dashed line shows the best-fit relation for individual H~II regions \\citet{cal05}.  The rest--frame 24~\\micron\\ luminosity for \\target\\ is significantly lower (by $\\approx 0.3-0.5$~dex) with respect to local galaxies of similar \\paa\\ luminosity.  In contrast, the $L(24\\micron)/L(\\paa)$ for target is consistent with that found in individual star--forming regions.   \\label{fig:24um}} \\end{figure*}  The observations for \\target\\ provide independent estimates for the total SFR.     In particular, we compare the \\paa\\ luminosity which stems from the ionized gas in \\ion{H}{2} regions and traces the number of ionizing photons, to estimates from the mid--IR and total IR luminosities, which measure primarily the dust--reprocessed emission from massive stars \\citep[\\eg,][]{ken98,kew02}.  As discussed in \\S~4.1, the \\paa\\ line luminosity corresponds to a SFR $\\psi = 171\\pm 28$~\\msol\\ yr$^{-1}$.  \\subsection{The \\paa\\ Luminosity Compared to the Total IR Luminosity}  Local star--forming galaxies and star--forming regions show a tight correlation between $L(\\paa)$ and \\lir\\  \\citep[\\eg,][]{cal05,alo06}. For \\target\\, R08 derived a total IR luminosity in the range $\\lir=5.7-9.5\\times 10^{11}$~\\lsol, taking  into account the systematic uncertainty owing to the choice of IR SEDs \\citep[see also the discussion in ][]{pap07}. Here, we reanalyze the total IR luminosity of \\target\\ by fitting different sets of IR SEDs to the flux density at 70, 450, and 850~\\micron\\ flux densities, all of which sample the thermal IR emission (see figure~\\ref{fig:irsed} and \\S~4.3).  We exclude the 24~\\micron\\ flux density from this analysis because it probes $\\sim$6--7~\\micron\\ in the rest--frame mid--IR, and its relationship to the thermal far--IR emission is not straightforward.  Figure~\\ref{fig:irsed} shows the best fit IR SED templates and  measured flux densities.   Using the \\citet{dal02} templates, we derive $\\lir = (5.0 \\pm 0.6) \\times 10^{11}$~\\lsol, with a goodness-of-fit, $\\chi^2/\\nu = 3.2$ and $\\nu=2$.  We find formally better fits using both the IR SEDs of \\citet{cha01}, giving $\\lir= (5.9\\pm 0.6) \\times 10^{11}$~\\lsol\\ with $\\chi^2/\\nu = 0.8$, and \\citet{rie09}, giving $\\lir = (10 \\pm 1.1)\\times 10^{11}$~\\lsol\\ with $\\chi^2/\\nu = 0.2$.  The IR luminosities from the fits are consistent with the results from R08, where the difference here is that we have excluded the 24~\\micron\\ flux density and include the deeper 70~\\micron\\ flux--density measurement.  Because the different IR SED templates all are consistent with the data, they imply there is a factor of 2 uncertainty on the total IR luminosity owing to differences in the choice of template.    Moreover, because the IR SEDs we have tested are not continuous in $\\lir$, this implies there is \\textit{at least} a factor of 2 uncertainty on total IR luminosities of $z\\sim 2.5$ galaxies, even when flux densities are available at 70~\\micron\\ and sub-mm wavelengths.  The total IR luminosities correspond to a range of SFR, $\\psi = 90-180$~\\msol\\ yr.   The upper range of the SFR corresponds to the fit using the Rieke et al.\\ template, and these are consistent with the SFR derived from the \\paa\\ luminosity.   The templates from Dale \\& Helou and Chary \\& Elbaz yield lower \\lir\\ values, and imply SFRs lower by $\\sim$40--50\\% compared to that from $L(\\paa)$.  To improve the accuracy on the IR luminosity will require flux density measurements at observed wavelengths of $\\approx$150--250~\\micron, in order to constrain the peak of the thermal dust emission (see figure~\\ref{fig:irsed}).  Nevertheless, the current analysis provides evidence that the total IR luminosity and \\paa\\ luminosities are consistent for \\target\\ within a factor of 2.  Parenthetically, we note that \\textit{none} of the empirical IR templates are able simultaneously to fit both the thermal IR emission and the strength of the PAH emission features in the mid--IR.    This is apparent in the lower panel of figure~\\ref{fig:irsed} where the best--fit IR templates imply lower observed 24~\\micron\\ flux densities less than the measured value.  This issue has been identified in the analysis of other  high--redshift gravitationally lensed galaxies \\citep[\\eg,][]{sia08}.  \\subsection{The \\paa\\ Luminosity Compared to the 8~\\micron\\ Luminosity}  The origin of the mid--IR emission (rest--frame 5--10~\\micron) in local galaxies is attributed both to very small-grain dust continuum emission and  molecular PAH emission.  Both the PAHs and very small grains are heated both by ionizing and non--ionizing sources (in particular, the ambient galactic radiation field from evolved stars, Li \\& Draine 2002).    Local galaxies and star--forming regions show a nonlinear relation between  $L(8\\micron)$ and $L(\\paa)$ \\citep[see, \\eg,][]{for04,cal07}, where the slope of the correlation depends on the dust--heating source(s), the dust spallation/formation rates, and the metallicity \\citep[\\eg,][]{hou04b,eng05,cal07,draine07}.  The MIPS 24~\\micron\\ flux density probes rest--frame mid--IR, $\\sim$6--7~\\micron\\ at $z$=2.515, as illustrated in figure~\\ref{fig:irsll}.  Many studies of the IR emission in distant galaxies use the measured 24~\\micron\\ flux density to estimate the total IR emission, and we test these relations for \\target.   Using the 24~\\micron\\ flux density with the prescription of \\citet{pap06} yields a estimated total IR luminosity, $\\lir = (1.2\\pm 0.1) \\times 10^{12}$~\\lsol, where the error is statistical only and does not include systematic uncertainties (see discussion in Papovich et al.\\ 2006).    This corresponds to a SFR, $\\psi \\simeq 220$~\\msol\\ yr$^{-1}$.  While this is consistent with the SFRs derived from the \\paa\\ emission and \\lir\\ measured from the far--IR data, this is somewhat fortuitous because the IR template used to extrapolate the observed 24~\\micron\\ flux density implies higher flux densities at observed--frame 70 and 850~\\micron\\ compared to the observations.  This is similar to the statement in \\S5.1 that \\textit{none} of the template IR SEDs are capable to fit simultaneously the far--IR flux densities and mid--IR emission features.   Using the scaling relation from \\citet{pap07}, which includes bolometric corrections using the average 70 and 160~\\micron\\ flux densities of $1.5 < z < 2.5$ galaxies, yields a nearly equal estimate for $\\lir$ as using the 24~\\micron\\ data only, with a similar offset compared to the implied SFR from \\paa.  R08 derive a scaling relation between the rest--frame 8~\\micron\\ luminosity, $L(8\\micron)$, and $\\lir$ for their sample of gravitationally lensed $z\\sim 2$ galaxies, which includes \\target. Using the $L(8\\micron)$ derived in \\S~4.3, the R08 scaling relation yields $\\lir = (7.9\\pm 1.5) \\times 10^{11}$~\\lsol. This is consistent within the range of the total IR luminosity derived above, and consistent with the implied SFR from \\paa\\ within the errors.  \\citet{pop08} derive scaling relations between \\lir\\ and the luminosity of the 6.2 and 7.7~\\micron\\ PAH emission features. R08 measured their PAH luminosities for \\target\\ simultaneously using PAHFIT (Smith et al.\\ 2007).  However, as discussed by Pope et al., line fluxes from measured by PAHFIT are generally higher than those using their method (see also the discussion in Sajina et al.\\ 2007; Siana et al.\\ 2008).   Thus, we remeasured the 6.2 and 7.7~\\micron\\ PAH luminosities individually, fitting each emission line with a Drude profile while fitting the slope and intercept of the continuum.   Our fits for each line yielded $L(6.2\\micron,\\mathrm{PAH}) = (2.25 \\pm 0.08)\\times 10^{43}$~erg s$^{-1}$ and $L(7.7\\micron, \\mathrm{PAH}) = (8.89 \\pm 0.04)\\times 10^{43}$~erg s$^{-1}$. Using equations 4 and 5 from Pope et al.\\ we infer $\\lir \\approx 2.3\\times 10^{12}$~\\lsol\\ and $3.1\\times 10^{12}$~\\lsol\\ for the 6.2 and 7.7~\\micron\\ PAH features, respectively.   While these estimates of \\lir\\ are higher by a factor of order three compared to that measured from the far-IR data,  they are within the scatter observed in the the relation between the PAH luminosities and total IR luminosity in the Pope et al.\\ sample. However, the intrinsic IR luminosity of \\target\\ is a factor of two lower than the high--redshift sub--mm from Pope et al., and it is possible the extrapolated relation does not hold.  This may be supported by the results of \\citet{sia08,sia09}, who observe a similar offset between the PAH luminosity and total IR luminosity in their study of intrinsically less-luminous, lensed UV--bright objects.  Figure~\\ref{fig:8um} shows the \\paa\\ luminosity against the rest--frame 8~\\micron\\ luminosity for \\target, compared  against luminosities for samples of local galaxies and  star-forming regions \\citep{cal05,cal07,eng05,dia08}.    \\citet{cal07} derive a scaling relation between the \\paa\\ and 8~\\micron\\ luminosities, $L(8\\micron) \\propto L(\\paa)^{\\alpha}$, with $\\alpha=0.94$. Note that Calzetti et al.\\ derived this correlation in terms of luminosity \\textit{surface densities} (luminosity per unit physical area).  However, \\citet{dia08} obtain a similar slope for the correlation between the 8~\\micron\\ and \\paa\\ \\textit{luminosities}.  While the luminosities for \\target\\ appear broadly consistent with the extrapolated relations,  a clearer picture is evident by comparing the ratio of $L(8\\micron)/L(\\paa)$, shown in the right panel of Fig~\\ref{fig:8um}.    There is much scatter in the local samples, but $L(8\\micron)/L(\\paa)$ for \\target\\ agrees broadly with those extrapolated relationships that include the star--forming galaxies \\citep{cal07,dia08}.  Using only the extrapolated relation from individual \\ion{H}{2} regions would underpredict the amount of 8~\\micron\\ emission \\citep[cf.,][]{cal05}.  This implies that  the fraction of photons from star--formation reradiated as 8~\\micron\\ luminosity is weakly dependent on ionizing luminosity.  \\subsection{The \\paa\\ Luminosity Compared to the 24~\\micron\\ Luminosity}  The emission at rest--frame 24~\\micron\\ results from thermal dust grains heated by ionizing and non--ionizing sources.  Empirically, local galaxies and star-forming regions follow a correlation with $L(24\\micron) \\propto L(\\paa)^\\alpha$, with $\\alpha$ in the range 1.03--1.23 \\citep{cal05,cal07,alo06}.  \\citet{alo06} argue that the correlation with $\\alpha > 1$ arises as dust absorbs ionizing and UV--continuum photons with increased efficiency in more heavily obscured, more luminous systems, so that an increasing fraction of the bolometric luminosity associated with star formation emerges in the IR with warmer dust temperatures \\citep{lon87,wang96,draineli07}.  A similar conclusion is reached by \\citet{cal07}, who argued that the $\\alpha > 1$ correlation exists because objects with increasingly higher starlight intensity and $L(\\paa)$ have higher dust temperatures, where the peak of the emission moves to shorter IR wavelengths.  \\citet{cal07} and \\citet{draineli07} discuss the physical basis for this correlation.   Figure~\\ref{fig:24um} shows the \\paa\\ luminosity plotted against the rest--frame 24~\\micron\\ luminosity derived above for \\target\\ and for local samples \\citep{cal05,cal07,eng05,dan05,alo06}.  Interestingly, \\target\\ has lower $L(24\\micron)$ than local galaxies of comparable $L(\\paa)$, where the offset is  $\\approx 0.3-0.5$~dex. This is more apparent when comparing the ratio of $L(24\\micron)/L(\\paa)$, shown in the right panel of figure~\\ref{fig:24um}.  Furthermore, because the 70~\\micron\\ data for \\target\\ is blended (see \\S~3.2), the rest--frame 24~\\micron\\ luminosity may be lower than indicated in the figure, making the offset between \\target\\ from the relation for the local samples more pronounced.  \\target\\ has an estimated total IR luminosity, $L_\\mathrm{IR}=(5-10)\\times 10^{11}$~\\lsol, comparable to local ultraluminous IR galaxies (ULIRGs, $\\lir \\ge 10^{12}$~\\lsol). Figure~\\ref{fig:24um} includes data for local ULIRGs from the sample of \\citet{dan05}, where we combine their \\paa\\ measurements with IRAS 25~\\micron\\ measurements from the literature\\footnote{see Moshir, Kopman, \\& Conrow (1992); the NASA Extragalactic Database (NED), http://nedwww.ipac.caltech.edu/}.  The Dannerbauer et al.\\ measurements of \\paa\\ come from longslit near--IR spectroscopy, and we have made no attempt to correct for emission outside the slit. The \\paa-emission in many of local ULIRGs likely results from very compact, nuclear regions, and therefore the spectroscopic slit should contain most of this emission. Nevertheless, we removed from the local ULIRG sample those objects with 2MASS isophotal diameters $>$15\\arcsec, and we also removed those objects with spectroscopic signatures of AGN.  These steps excluded roughly one--third of the sample, including the most egregious outliers. Nevertheless we caution that significant uncertainty may remain due to primarily unknown \\paa\\ emission outside the spectroscopic slit or other aperture effects.  Even with these caveats the local ULIRGs follow the extrapolation of the local $L(24\\micron)$ and $L(\\paa)$ relation, but they show a large scatter.  We suspect that the large scatter results for the reasons discussed above, and there may be additional components to the dust heating beyond ionization from early--type stars, including AGN and the ambient galactic emission. However, only one galaxy in the local ULIRG sample has $L(24\\micron)/L(\\paa)$ and $L(\\paa)$ comparable to \\target, and this galaxy (IRAS 04384--4848) has a highly uncertain dust correction \\citep{dan05}.  This implies that no (or at best, \\textit{few}) low redshift ULIRGs have similar physical conditions producing comparable ratios of mid-IR-to-\\paa\\  luminosity.  \\citet{ken09} combine \\ha\\ emission--line measurements (uncorrected for dust extinction) and IR continuum measurements of local star--forming galaxies, and derive SFR calibrations of the form, $\\psi = 7.9\\times 10^{-42} \\times [ L(\\ha)_\\mathrm{obs} + a_\\lambda L(\\lambda)]$. For the IRAC 8~\\micron\\ and MIPS 24~\\micron\\ rest--frame bands they obtain $a_{8} = 0.011$ and $a_{24}=0.020$.   Using the mid--IR luminosities for \\target\\ derived above, and $L(\\ha)_\\mathrm{obs} = 3.7 \\times 10^{8}$~\\lsol\\ \\citep{kne04}, we obtain $\\psi \\simeq 45$ and 80~\\msol\\ yr$^{-1}$, for the 8~\\micron\\ and 24~\\micron\\ luminosities, respectively.   These are lower by factors of 4 and 2 compared to that derived from the \\paa\\ luminosity, but they are within the dispersion reported by \\citet{ken09}.  The intrinsic SFR of \\target\\ is also considerably larger than the objects used to calibrate these relations (see also Moustakas \\& Kennicutt 2006), and it is possible the calibrations do not apply under extrapolation.  Larger samples of luminous, high--redshift galaxies are needed to test these relations.  \\begin{figure}[t] \\ifsubmode \\epsscale{1.0} \\else \\epsscale{1.2} \\fi \\plotone{f8.eps} \\epsscale{1.0} \\caption{ Infrared SED of \\target\\ compared to radiative transfer models of \\citet{sie07}.  The data points show the measured flux densities at IRAC 3.6, 4.5, 5.8 and 8.0~\\micron, the MIPS 24 and 70~\\micron, and the SCUBA 450 and 850~\\micron.  The right axis shows the flux densities after correcting for the gravitational lensing magnification.  The top axis shows the rest--frame wavelength for $z=2.515$.  The curves show model fits to the 24, 70, 450, and 850~\\micron\\ flux densities.  These SEDs are computed for a spherical PDR ionized uniformly by an interior starburst.  The resulting amount of visual extinction is a variable in the model.  Each curve shows the best--fit model as a function of visual extinction, as indicated in the figure inset. \\label{fig:irrs07}} \\end{figure}  \\subsection{The Nature of \\target}  \\target\\  appears to host heavily obscured star--formation at a rate, $\\psi \\approx 170$~\\msol\\ yr$^{-1}$.   The hydrogen ionization rate is $Q_0 = (1.6\\pm 0.3) \\times 10^{55}$~$\\gamma$ s$^{-1}$, implying \\target\\ may contain as many of $\\sim 10^6$ O stars \\citep{stern03}.   Given  the estimate for the molecular--gas mass ($M[\\mathrm{H}_2] \\approx 4.5\\times 10^9$~\\msol, Kneib et al.\\ 2005),  this galaxy could sustain this SFR for $t \\sim 30$~Myr.   The starburst in \\target\\ has had a duration of $\\le$100 Myr  based on the analysis of rest--frame UV-to-near-IR SED (see the Appendix) and supported by the strength of the possible CO index (\\S~4.2).  This is comparable to the dynamical and gas--consumption timescales based on the observations of the molecular gas.   Because the gas--consumption timescale is consistent with the dynamical time and the starburst age, \\target\\ is likely about midway through this stage of enhanced star formation.  The \\paa--derived SFR for \\target\\ is consistent with the SFR implied by the  total IR luminosity and the rest--frame 8~\\micron\\ luminosity. However, the rest--frame 24~\\micron\\ luminosity is significantly lower than that expected from the \\paa\\ luminosity based on the local relations.    This implies that  \\target\\ lacks a warm ($\\sim$100~K) thermal dust component typical of local \\textit{IRAS}--selected star--forming galaxies of comparable bolometric luminosity \\citep{lon87,cal00}, which drives the non--linear relationship between the mid--IR luminosity and the \\paa\\ luminosity \\citep[\\eg,][]{cal07}. Indeed, the thermal dust emission of \\target\\ (see figure~\\ref{fig:irsed}) peaks near $\\sim$100~\\micron, and is consistent with nearly pure emission from dust with temperature, $T_D = 52$~K, and emissivity, $\\beta=1.5$.  This is very similar to the dust emission of photodissociation regions (PDRs) in the vicinity of \\ion{H}{2} regions \\citep[\\eg,][]{lon87,cal00,chu02}.  In contrast, the IR SEDs for local galaxies with $\\lir$ \\textit{comparable to } \\target\\ have significant contributions of warm ($\\gsim70$~K) dust to the IR emission \\citep[\\eg,][]{cha01,dal02,sie07,rie09}.  This is a qualitatively different than what is observed in \\target.  Indeed, the fact that no galaxies in the local ULIRG sample have $L(24\\micron)  / L(\\paa)$ values as low as \\target\\ implies that star formation in some high--redshift galaxies with ULIRG luminosities is fundamentally different. The lower ratio, $L(24\\micron) / \\lir \\simeq 0.1$, for \\target\\ implies a lower starlight intensity than that in local galaxies of comparable bolometric luminosity \\citep{draineli07}. Therefore, while \\target\\ appears to host a massive starburst that is similar to local star--forming regions, it is dramatically ``scaled--up'' in luminosity (and presumably SFR).  This conclusion is supported by a comparison of the mid--IR and far--IR emission of \\target\\ to the models of \\citet{sie07}, who calculate the IR emission for PDR--like regions with spherical symmetry of variable size surrounding \\ion{H}{2} regions ionized by a starburst with a variable fraction of the luminosity coming from OB associations.   As illustrated in figure~\\ref{fig:irrs07}, the best--fitting models from Siebenmorgen \\& Kr\\\"ugel correspond to an obscured stellar population where $60-90$\\% of the luminosity originates from OB associations.   Such models require emission from dust clouds with a range of extinction, $A(V) \\approx 2-72$~mag. Models with visual extinction $A(V) \\leq 18$~mag that reproduce the data have an \\textit{intrinsic} luminosity $\\approx10^{10}-10^{11}$~\\lsol, and this must be scaled up by a factors of $\\approx$10--100 to match the data, producing total IR luminosities in agreement with that derived above (\\S~5.1).  Scaling the size of the PDR according, this implies a radius for the ionization--front of $\\sim$3-4~kpc, consistent with the spatial extent of the molecular CO emission of \\target\\ \\citep[$\\sim$3~kpc $\\times$ 1.5~kpc,][]{kneib05}.  Models with higher \\textit{intrinsic} IR luminosity ($\\gg10^{11}$~\\lsol) do not reproduce the data.   Models with higher extinction require larger intrinsic luminosities (as much as an order of magnitude larger for the model with $A(V) = 72$~mag), and these provide worse fits to the data (thus, they are less physical).  Furthermore, models with $A(V) \\gsim 30$~mag would be optically thick to \\paa\\ photons, and cause significant attenuation of the mid--IR emission.   While such models are physically motivated in some cases \\citep[\\eg,][]{rie09}, if this were the case for \\target\\, then the attenuation correction to the \\paa\\ luminosity would imply a much larger SFR, which should be substantiated by a larger total IR luminosity.  The general agreement between the dust--corrected \\paa\\ luminosity and total IR luminosity (\\S5.1) excludes models with very high extinction.  While the models assume a spherical shell-like dust configuration, in reality the dust clouds are likely clumpy with some covering fraction, where the ionizing \\ion{H}{2} regions associated with the different OB associations are in close proximity to the dust clouds and PDRs \\citep{wol90}.  Models factoring in the covering fraction produce a distribution of dust attenuation \\citep[\\eg,][]{dop05}, and this is more consistent with the observations.    Figure~\\ref{fig:irrs07} shows that IRAC 3.6, 4.5, 5.8 and 8.0~\\micron\\ flux densities for \\target\\ are not reproduced simultaneously with the far--IR emission by any of the best-fitting models.  Thus, spherical symmetry of the starburst and obscuring PDR seems insufficient to describe both the direct stellar emission and the far--IR emission for this galaxy. Therefore, we conclude that star formation in \\target\\ corresponds to star--forming regions and starbursts in local galaxies with intrinsic luminosities of $\\lir \\approx 10^{10}-10^{11}$~\\lsol, but that have been ``scaled--up'' by one--to--two orders of magnitude.   Moreover, this ``scaling--up'' of star formation may be common at high redshift.  For example, \\citet{tac06} argued that the significantly more luminous sub-mm galaxies at $z\\sim 2-3$ resembled ``scaled--up'' and more gas--rich versions of local ULIRGs.   For \\target,  R08 found that the rest--frame mid--IR spectrum from $6-10$~\\micron\\ is consistent with the spectra of local starburst galaxies and inconsistent with the spectra of local ULIRGs.  Here, our analysis of the far--IR data, coupled with the SFR implied by the luminosity of the \\paa\\ line, leads us to a similar conclusion.  The model we constructed in the Appendix using the observed rest--frame UV--near-IR SED  consists of a double stellar population. The stellar population that dominates the bolometric emission (and the total SFR) in \\target\\ is very obscured by dust that is optically thick to UV photons.  The stellar population of the subdominant model is less obscured, and optically thin to UV photons.   This analysis is consistent with the scenario discussed above, where the starburst in \\target\\ corresponds to OB associations in close proximity to PDRs with a clumpy distribution.   A similar configuration of multiple star--forming components with variable dust extinction is observed in the spatially resolved colors of other galaxies at $z\\sim 2-3$, and especially those with evidence for recent mergers \\citep{pap05,law07}.  It remains to be seen whether the results discussed here for \\target\\ are typical of other high--redshift galaxies.  Our full observing program will produce similar data for eleven gravitationally lensed galaxies in addition to \\target, and these will span a wide range of mass, optical, and IR properties.  For \\target, the baryon (stellar + gas) mass we derived from the rest-frame-UV--to--near-IR SED and the dynamical mass from molecular observations \\citep{kneib05} both suggest $M \\sim 1.5 \\times 10^{10}$~\\msol.  This mass is typical of ``$L^\\ast$'' UV--selected galaxies at these redshifts \\citep{pap01,sha01}, and such objects likely dominate the SFR density at this redshift \\citep{red08}.    However, the dust obscuration and IR luminosity are much larger in \\target\\ compared to the UV--selected samples, and more typical of sub-mm--selected objects \\citep[\\eg][]{blain02,chap05} and IR--luminous $K$--band-selected objects \\citep[\\eg,][]{pap06,dad07,wuyts08}.  The fact that the molecular--gas--emission centroid lies between the UV--components led \\citet{kneib05} to argue that the high SFR in \\target\\ results from the merger of two progenitors whose properties are similar to those of ``typical'' UV--selected galaxies.     In this case, we are observing \\target\\ during perhaps a short--lived stage of enhanced star formation.   Our study of the SFR indicators in \\target\\ provides the first evidence that the total IR luminosity is proportional to the number of ionizing photons in these situations.  \\vspace{12pt}  \\subsection{Final Thoughts}  A potential source of bias is that our observations are sensitive to the integrated emission from \\target, and we are combining datasets with a wide range of angular resolution.     This is of concern as as our analysis corresponds to the flux--weighted average of the star--formation properties of an individual galaxy.   The analysis of the rest--frame UV--to--near-IR SED (see the Appendix) suggests there are \\textit{at least} two stellar populations, with very different extinction properties.    Nevertheless, our results probably apply to the dominant, most luminous star--forming components, which dominate the nebular emission (\\paa) and the IR emission.   Therefore, we suspect that our results are valid for the average properties of \\target.   We note that similar issues arise in the analysis of the integrated properties of the starburst galaxies \\citep{eng05,alo06,cal07}, and some caution must be applied then when comparing these with, for example, the analysis of individual \\ion{H}{2} regions which are resolved in both the narrow--band imaging of the hydrogen recombination lines and the mid--IR data \\citep{cal05,cal07}.  Achieving resolved observations of the \\paa\\ emission and far--IR emission in \\target\\ may be possible with the \\textit{James Webb Space Telescope} and large--aperture sub--mm facilities.   However, further work will be needed to quantify possible biases in the integrated emission from galaxies."
    },
    "0910/0910.0498_arXiv.txt": {
        "abstract": "The $^{14}$C($\\alpha$,$\\gamma$) reaction rate at temperatures below 0.3 GK depends on the properties of two near threshold resonances in $^{18}$O, the 1$^{-}$ at 6.198 MeV and the 3$^{-}$ at 6.404 MeV. The $\\alpha$+$^{14}$C Asymptotic Normalization Coefficients (ANCs) for these resonances were determined using the $\\alpha$-transfer reactions $^{14}$C($^{7}$Li,t) and $^{14}$C($^{6}$Li,d) at sub-Coulomb energies. The $^{14}$C($\\alpha$,$\\gamma$) reaction rate at low temperatures has been evaluated. Implications of the new reaction rate on the evolution of accreting helium white dwarfs and on the nucleosynthesis of low mass stars during the asymptotic giant branch (AGB) phase are discussed. ",
        "introduction": "}  The $^{14}$C($\\alpha$,$\\gamma$)$^{18}$O reaction has been suggested to play an important role in several astrophysical environments. It was pointed out by Mitalas \\cite{Mitalas1974} that electron capture on $^{14}$N can produce $^{14}$C in the degenerate helium core of low mass stars. The $\\alpha$-capture on $^{14}$C then follows. It was suggested in \\cite{Kaminisi1975} that the $^{14}$N(e$^{-}$,$\\nu$)$^{14}$C($\\alpha$,$\\gamma$)$^{18}$O reaction (NCO reaction) may trigger the helium flash in the core of low mass stars earlier than the $3\\alpha$ reaction. The significance of the NCO reaction for the evolution of low mass stars was also considered by Spulak \\cite{Spulak1980} who concluded that it is not effective compared to the 3$\\alpha$ reaction. A contradicting result was obtained by Kaminisi and Arai \\cite{Kaminisi1983} who used the updated (but still very uncertain) $^{14}$C($\\alpha$,$\\gamma$) reaction rate and concluded that the NCO reaction is dominant for igniting the helium flash. While the influence of the NCO reaction on the evolution of low mass stars in the red giant phase is still uncertain it was firmly established by Nomoto and Sugimoto \\cite{Nomoto1977}, and later by Hashimoto, et al. \\cite{Hashimoto1986}, that the NCO reaction may trigger the helium flash in accreting helium white dwarfs at a lower temperature and density than the $3\\alpha$ reaction. However, this conclusion is rather sensitive to the actual value of the $^{14}$C($\\alpha$,$\\gamma$) reaction rate at temperatures between 0.03 and 0.1 GK.  The $^{14}$C($\\alpha$,$\\gamma$) reaction is also important for production of $^{19}$F in asymptotic giant branch stars. Observations by Jorissen, et al., \\cite{Jorissen1992} show enhanced fluorine abundance in the atmosphere of K, M, MS, S, SC, and C asymptotic giant branch (AGB) stars. This finding indicates that $^{19}$F is produced in these stars and its abundance can be used to constrain the properties of the AGB models. The path for $^{19}$F production in an AGB star, as proposed by Jorissen \\cite{Jorissen1992}, is rather complex, occurring in the He intershell of the AGB star. Neutrons from the s-process neutron generator reaction, $^{13}$C($\\alpha$,n), can be captured by $^{14}$N, enhanced from the previous CNO cycle. This leads to production of $^{14}$C and protons through the $^{14}$N(n,p)$^{14}$C reaction. Then $^{18}$O is produced by the $^{14}$C($\\alpha$,$\\gamma$) reaction, or, alternatively by $^{14}$N($\\alpha$,$\\gamma$)$^{18}$F($\\beta$)$^{18}$O. The presence of protons and $^{18}$O isotopes in the He intershell simultaneously allows for the $^{18}$O(p,$\\alpha$)$^{15}$N reaction to occur and subsequently $^{15}$N($\\alpha$,$\\gamma$)$^{19}$F capture. Competing reactions that reduce the final abundance of $^{19}$F are $^{19}$F($\\alpha$,p), $^{18}$O($\\alpha$,$\\gamma$)$^{22}$Ne and $^{15}$N(p,$\\alpha$)$ $$^{12}$C. A detailed investigation of the $^{19}$F production in AGB stars has been performed recently by Lugaro, et al., \\cite{Lugaro2004} and it was determined that the major uncertainties in the production of $^{19}$F are associated with the uncertainties in the $^{14}$C($\\alpha$,$\\gamma$)$^{18}$O and $^{19}$F($\\alpha$,p)$^{22}$Ne reaction rates. Addressing uncertainties of the $^{14}$C($\\alpha$,$\\gamma$) reaction rate at temperatures relevant for helium flashes in the accreting helium white dwarfs and nucleosynthesis in the asymptotic giant branch stars is the main goal of this work.  \\begin{figure} \\includegraphics[width=0.9\\columnwidth]{18OLevels.eps}  \\caption{\\label{fig:18Olevels}Partial level scheme of $^{18}$O.}  \\end{figure}  Direct measurements of the $^{14}$C($\\alpha$,$\\gamma$) reaction cross section are only available for energies above 880 keV in c.m. \\cite{Gorres1992}. The relevant temperature range for accreting helium white dwarfs is between 0.03 and 0.1 GK and for $^{19}$F nucleosynthesis in AGB stars it is $\\sim$0.1 GK. The Gamow window for the $^{14}$C($\\alpha$,$\\gamma$) reaction at these temperatures corresponds to an energy range between 50 and 250 keV. At these energies the cross section for the $\\alpha$ capture reaction on $^{14}$C is too low to be measured directly and has to be extrapolated from the higher energy data. However, near threshold resonances in $^{18}$O may have a very strong influence on this extrapolation. $^{}$The level scheme of $^{18}$O is shown in Fig. \\ref{fig:18Olevels}. Three states with excitation energies close to the $\\alpha$ decay threshold at 6.227 MeV are known in $^{18}$O. These are the 1$^{-}$ at 6.198 MeV, the 2$^{-}$ at 6.351 MeV and the 3$^{-}$ at 6.404 MeV. The 2$^{-}$ state at 6.351 MeV is an unnatural parity state and cannot contribute to the $^{14}$C($\\alpha$,$\\gamma$) reaction at any significant level. The 1$^{-}$ subthreshold state at 6.198 MeV can only contribute to the $^{14}$C($\\alpha$,$\\gamma$) reaction through its high energy tail, therefore its contribution is expected to be relevant only at the lowest energy. Finally, the 3$^{-}$ state at 6.404 MeV is 177 keV above the $\\alpha$-decay threshold - right in the middle of the energy range of interest.  The strength of a resonance is defined by\\begin{equation} \\omega\\gamma=\\frac{2J_{R}+1}{(2j_{t}+1)(2j_{p}+1)}\\frac{\\Gamma_{\\alpha}\\Gamma_{\\gamma}}{\\Gamma},\\end{equation}   \\noindent where $\\Gamma$, $\\Gamma_{\\alpha}$ and $\\Gamma_{\\gamma}$ are the total, partial $\\alpha$ and partial $\\gamma$ widths and $J_{R}$, $j_{t}$ and $j_{p}$ are spins of the resonance, the target and the projectile. At 177 keV the partial $\\alpha$-width is much smaller than the partial $\\gamma$ width ($\\Gamma_{\\alpha}\\ll\\Gamma_{\\gamma}$), therefore the resonant $\\alpha$-capture cross section is determined entirely by the partial $\\alpha$-width with resonance strength calculated as $\\omega\\gamma\\approx7\\times\\Gamma_{\\alpha}$ for the 3$^{-}$ state at 6.404 MeV in $^{18}$O. The partial $\\alpha$-width of the 3$^{-}$ state is not known. This leads to several orders of magnitude uncertainty of the $^{14}$C($\\alpha$,$\\gamma$) reaction rate at $\\sim0.1$ GK. The subthreshold 1$^{-}$ state at 6.198 MeV also contributes to the reaction rate uncertainty at the lowest temperatures.  Information on the partial $\\alpha$ widths of resonances in $^{18}$O can be extracted using the direct $\\alpha$-transfer reactions $^{14}$C($^{6}$Li,d) and $^{14}$C($^{7}$Li,t). Usually these reactions are performed at energies of 5-15 MeV/A and $\\alpha$ spectroscopic factors (S$_{\\alpha}$) for the states of interest are determined from the Distorted Wave Born Approximation (DWBA) analysis of the reaction cross sections. The partial $\\alpha$ width is then calculated as $\\Gamma_{\\alpha}=S_{\\alpha}\\times\\Gamma_{SP}$, where $\\Gamma_{SP}$ is the $\\alpha$ single particle width.  \\begin{eqnarray} \\Gamma_{SP} & = & 2P_{\\ell}(kR)\\gamma_{sp}^{2}\\label{eq:width}\\\\ P_{\\ell}(kR) & = & \\frac{kR}{F_{\\ell}^{2}(k,R)+G_{\\ell}^{2}(k,R)}\\label{eq:pen}\\\\ \\gamma_{sp}^{2} & = & \\hbar^{2}/\\mu R^{2},\\label{eq:redwidth}\\end{eqnarray}   \\noindent where $P_{\\ell}(kR)$ is a penetrability factor determined by the Coulomb regular and irregular functions $F_{\\ell}(k,R)$ and $G_{\\ell}(k,R)$, $\\mu$ is a reduced mass and $k$ is a wave number. R is a channel radius - $R=r_{0}(\\sqrt[3]{A_{1}}+\\sqrt[3]{A_{2}})$, where $r_{0}=1.2-1.4$ fm. The reduced width, $\\gamma_{sp}$, in equation \\eqref{eq:redwidth} corresponds to the reduced width in a square-well potential. The single particle width can be calculated more accurately using a realistic Woods-Saxon potential. As is well known, the spectroscopic factor extracted from an $\\alpha$-transfer reaction depends on the parameters of the optical potentials used to describe the wave functions of relative motion in the entrance and exit channels in the DWBA approach and on the shape of the form factor potentials and the number of nodes in the model core-cluster wave function. However, for astrophysical calculations the asymptotic normalization coefficient (ANC) is needed instead of the spectroscopic factor, and if the transfer reaction is performed at sub-Coulomb energies, then the parametric dependence of the DWBA calculations is greatly reduced. This approach has been applied before in Refs. \\citet{Brune1999,Johnson2006} where the $^{12}$C($\\alpha$,$\\gamma$) and $^{13}$C($\\alpha$,n) reaction rates due to near threshold resonances in $^{16}$O and $^{17}$O were evaluated. In this work we measured the cross sections of the $^{14}$C($^{7}$Li,t) and $^{14}$C($^{6}$Li,d) $\\alpha$-transfer reactions at sub-Coulomb energies and determined the ANCs of the 1$^{-}$ and 3$^{-}$ states in $^{18}$O. The discussion is structured as follows: the experimental technique is discussed in the section \\ref{sec:Experiment}, the extraction of the ANCs and associated uncertainties are described in section \\ref{sec:Analysis}, the new $^{14}$C($\\alpha$,$\\gamma$) reaction rate and it's astrophysical implications are discussed in section \\ref{sec:rate}, and the conclusions are in section \\ref{sec:Conclusions}. ",
        "conclusions": "}  The $^{14}$C($\\alpha$,$\\gamma$) reaction rate was studied using an indirect approach. The asymptotic normalization coefficients (ANCs) for the 1$^{-}$ state at 6.198 MeV and 3$^{-}$ at 6.404 MeV in $^{18}$O were measured in the $\\alpha$-transfer reactions ($^{7}$Li,t) and ($^{6}$Li,d) on $^{14}$C. The measurements were performed at sub-Coulomb energies to minimize the dependence of the final result on the optical model potential parameters. The extraction of ANCs instead of spectroscopic factors significantly reduced the dependence on the form-factor potential parameters and the number of nodes assumed in the $\\alpha$-$^{14}$C wave function. The resonance strength of the 3$^{-}$ state was determined from the ANC to be $\\omega\\gamma=(5.5\\pm1.9)\\times10^{-13}$ eV. It was found that the resonance capture due to the 3$^{-}$ state dominates the reaction rate in the temperature range of 0.03<T$_{9}$<0.3 which is the most relevant temperature range for the $^{19}$F nucleosynthesis in AGB stars, evolution of helium accreting white dwarfs and core helium flashes of low-mass stars. At temperatures below 0.03 GK the $^{14}$C($\\alpha$,$\\gamma$) reaction rate is determined by the interplay between the direct capture and capture due to the sub-threshold 1$^{-}$ resonance at 6.198 MeV. The S-factor due to the 1$^{-}$ sub-threshold state was determined in this work from the ANC and the direct capture S-factor was suggested in \\cite{Gorres1992}. In spite of this knowledge, the reaction rate at the temperatures below 0.03 GK is still uncertain by two orders of magnitude due to the unknown interference (constructive or destructive) between the two amplitudes.  The new $^{14}$C($\\alpha$,$\\gamma$) reaction rate is a factor of 30 lower than the one used in \\cite{Hashimoto1986}. Therefore, the importance of the NCO chain as a trigger for helium flashes in helium accreting white dwarfs suggested in \\cite{Hashimoto1986} is reduced, if not eliminated all together. The {}``recommended'' $^{14}$C($\\alpha$,$\\gamma$) reaction rate used in \\cite{Lugaro2004} for $^{19}$F nucleosynthesis calculation in AGB stars fortuitously coincides with that determined from our experimental data and the uncertainty of the $^{19}$F production in AGB stars associated with the $^{14}$C($\\alpha$,$\\gamma$) reaction rate is now eliminated. We hope that the new information on the $^{14}$C($\\alpha$,$\\gamma$) reaction rate will help to resolve the question of whether or not the NCO reaction chain can trigger helium flashes in cores of low-mass stars before the 3$\\alpha$ reaction does.  Authors are grateful to Peter Hoeflich and Akram Mukhamedzhanov for enlightening discussions and acknowledge support of National Science Foundation under grant number PHY-456463."
    },
    "0910/0910.0721_arXiv.txt": {
        "abstract": "We present 737 candidate Young Stellar Objects (YSOs) near the W51 Giant Molecular Cloud (GMC) over an area of $1.25 \\arcdeg \\times 1.00\\arcdeg $ selected from {\\it Spitzer Space Telescope} data. We use spectral energy distribution (SED) fits to identify YSOs and distinguish them from main-sequence or red giant stars, asymptotic giant branch stars, and background galaxies. Based on extinction of each YSO, we separate a total of 437 YSOs associated with the W51 region from the possible foreground sources. We identify 69 highly embedded Stage 0/I candidate YSOs in our field with masses ${>}~5 ~M_{\\odot}$ (corresponding to mid- to early-B main-sequence spectral types) 46 of which are located in the central active star forming regions of W51A and W51B. From the YSOs associated with W51, we find evidence for mass segregation showing that the most massive YSOs are concentrated on the W51 \\ion{H}{2} region complex. We find a variation in the spatial distribution of the mass function (MF) of YSOs in the mass range between 5 $M_\\odot$ and 18 $M_\\odot$. The derived slopes of the MF are $-1.26$ and $-2.36$ in the active star-forming region and the outer region, respectively. The variation of the MF for YSOs embedded in the molecular cloud implies that the distribution of stellar masses in clusters depends on the local conditions in the parent molecular cloud. ",
        "introduction": "Stars form in various environments. High-mass stars are especially important because they affect their environment through such phenomena as outflows, stellar winds, and supernovae. To understand the interaction of stars and the interstellar medium (ISM) it is necessary to identify embedded young stellar objects (YSOs) in the parent molecular cloud. Therefore taking a complete census of YSOs is an important step toward understanding YSOs and their environment on a large scale.  W51 is an active star-forming region located in the Sagittarius spiral arm. The distance is uncertain, with published values including $5.1^{+2.9}_{-1.4}$ kpc \\citep{Xu09}, $6.1\\pm1.3$ kpc \\citep{Imai02}, or $8.3\\pm2.5$ kpc \\citep{Schneps81}. The radio continuum sources comprising W51 are within a massive giant molecular cloud (GMC) \\citep{Mufson79}. Observations showing the star forming activity have been carried out in various wavelengths of radio continuum, HI line, molecular lines, water masers, near-IR, X-ray \\citep{Bieging75, Carpenter98, Imai02, Koo97a, Koo97, Koo99, Mehringer94, Okumura00}. W51 is a region of ongoing massive star formation \\citep{Lacy07, Zapata08} and therefore a good place to study the feedback between stars and the ISM.  To understand the interaction of molecular clouds and YSOs in the W51 \\ion{H}{2} region complex we mapped a 1.25\\arcdeg\\ $\\times$ 1.0\\arcdeg\\ region in the $J=2-1$ transitions of \\co and \\tco with the 10 m Heinrich Hertz Telescope (HHT) on Mount Graham, Arizona. The molecular line maps will be presented in a separate paper (Bieging et al., in preparation). In the present paper, we use {\\it Spitzer} data for identifying the YSO candidates associated with the molecular clouds. {\\it Spitzer} surveys in the mid-IR allow sensitive studies of deeply embedded star formation. Although extinction is significantly reduced in the mid-IR, it is still important for active star forming regions in giant molecular clouds. In particular, the W51 region shows no detectable increase in star counts because of severe extinction \\citep{Benjamin05}. Therefore we have to consider extinction effects carefully in identifying and characterizing candidate YSOs. Combining the YSOs selected from the {\\it Spitzer} photometry and molecular clouds with kinematic information, we can examine the feedback process between star(s) and ISM. The detailed analysis for the interaction of YSOs and associated molecular clouds will be presented in other publications (\\cite{Kang09}; Kang et al., in preparation). In this paper, we tabulate all of the YSO candidates and discuss their spatial distribution and mass function. In Section \\ref{sec:yso_selection}, we introduce the data sets used and describe our methods for selecting and classifying candidate YSOs. We discuss some of the implications of our results in Section \\ref{sec:result} and summarize the results in Section \\ref{sec:summary}. ",
        "conclusions": ""
    },
    "0910/0910.4102_arXiv.txt": {
        "abstract": "We present the results of a 8.4 GHz Very Large Array radio survey of early-type galaxies extracted from the sample selected by C{\\^o}t{\\'e} and collaborators for the Advanced Camera for Survey Virgo Cluster Survey. The aim of this survey is to investigate the origin of radio emission in early-type galaxies and its link with the host properties in an unexplored territory toward the lowest levels of both radio and optical luminosities. Radio images, available for all 63 galaxies with B$_{\\rm T} < 14.4$, show the presence of a compact radio source in 12 objects, with fluxes spanning from 0.13 to 2700 mJy. The remaining 51 galaxies, undetected at a flux limit of $\\sim$0.1 mJy, have radio luminosities L $\\lesssim 4 \\times 10^{18} \\, {\\rm W Hz}^{-1}$. The fraction of radio-detected galaxies are a strong function of stellar mass, in agreement with previous results: none of the 30 galaxies with M$_{\\star} < 1.7 \\times 10^{10} M_{\\sun}$ is detected, while 8 of the 11 most massive galaxies have radio cores. There appears to be no simple relation between the presence of a stellar nucleus and radio emission. In fact, we find radio sources associated with two nucleated galaxies, but the majority of nucleated objects are not a radio emitter above our detection threshold.  A multiwavelength analysis of the active galactic nucleus (AGN) emission, combining radio and X-ray data, confirms the link between optical surface brightness profile and radio loudness in the sense that the bright core galaxies are associated with radio-loud AGNs, while non-core galaxies host radio-quiet AGNs.  Not all radio-detected galaxies have a X-ray nuclear counter part (and vice-versa). A complete census of AGNs (and supermassive black holes, SMBHs) thus requires observations, at least, in both bands. Nonetheless, there are massive galaxies in the sample, expected to host a large SMBH (M$_{\\rm BH} \\sim 10^8 M_{\\sun}$), whose nuclear emission eludes detection despite their proximity and the depth and the spatial resolution of the available observations. Most likely this is due to an extremely low level of accretion onto the central SMBH. ",
        "introduction": "\\label{intro}  Nuclear radio emission is almost invariably associated with the presence of an active galactic nucleus (AGN) and this is an indication that the process of accretion onto a supermassive black hole (SMBH) naturally produces a signature in the form of radiation in the radio domain. The separation between radio-loud (RL) and radio-quiet (RQ) AGNs is in fact only a measure of the relative flux in the radio band with respect to the optical or X-ray nucleus \\citep{kellermann94,terashima03}; but also RQ AGNs, when studied at sufficient depth, usually show the presence at least of a nuclear radio component \\citep[e.g.,][]{ulvestad89,nagar05}. For RL AGNs the radio core results from synchrotron emission produced by the unresolved base of their jets; for RQ AGNs the situation is more controversial and it has been recently proposed, besides the possibility of a jet origin, that their radio nuclei are the manifestation of the presence of a thermal outflow or of an active disk corona \\citep{blundell07,laor08}  When combined to the very limited effects that absorption has on radio waves, the study of radio emission provides, in principle, a very powerful tool to detect accretion onto SMBH and, consequently, to establish when a SMBH is present in a given galaxy. However, to take full advantage of this approach, we must reach a much deeper understanding of what determines the radio luminosity of a given galaxy and how this is related to its level of accretion.  A large effort has been dedicated to explore the connection between the host properties (mostly from an optical point of view) and its radio emission. Already from the pioneering study by \\citet{auriemma77} it was clear that more massive galaxies have on average a higher radio luminosity than smaller galaxies, while apparently there are no distinctions between clusters and non-clusters members \\citep{ledlow96}. The radio luminosity functions (RLFs) of galaxies of different optical magnitudes are similar but they differ strongly in their scaling. More recent studies confirm the early results, indicating that the normalization of the RLF scales with the host luminosity as $\\sim$L$^{2.5}$ \\citep[e.g.,][]{best05b,mauch07}. However, galaxies of given optical magnitude show a very large range of radio power, more than 5 orders of magnitude, and the relation between the radio and optical luminosity can only be described in terms of a probability distribution.  These studies focused mostly on massive galaxies and had a relatively high threshold for the radio detection. The analysis reaching the lowest level of luminosity (in both bands) were performed by \\citet{sadler89} and \\citet{wrobel91b}, with limits at L$_{\\rm r} \\sim 2.5 \\times 10^{19} \\, {\\rm W Hz}^{-1}$ and M$_{\\rm r} \\sim $ -18.5. Consequently, the information on the radio properties of galaxies of lower mass is sparse and incomplete. Similarly, there is still a dominant fraction of massive galaxies ($\\sim 70$ \\%) undetected in radio surveys, for which only an upper limit to their radio luminosity can be derived.  Clearly, a survey reaching lower radio and optical luminosities can provide new insights on the origin and properties of the radio emission of early-type galaxies. The Virgo cluster represents a unique laboratory for such studies. In fact, it includes hundreds of early-type galaxies, spanning a wide range of stellar masses, for which a vast suite of data is available in the literature. In particular, the recent {\\it Hubble Space Telescope} ({\\it HST}) survey performed by \\citet{cote04} provides us with a detailed analysis of their optical brightness profiles whose properties, including the presence of a stellar nucleus, can be included in the study of their radio emission.  Given its proximity, its members can be observed at high spatial resolution (1$\\arcsec$ at a distance of 17 Mpc corresponds to $\\sim$ 80 pc) and their nuclear emission can be studied down to extremely low luminosity level. For this reason we performed a radio survey of 63 early-type galaxies in Virgo with the Very Large Array (VLA). Since the aim of these observations is to explore their nuclear properties, the data were taken at high resolution (with the telescope in the A array configuration) and at relatively high frequency (8.4 GHz).  The paper is organized as follows. In Section \\ref{sample} we describe the sample and the VLA observations; in Section \\ref{results} we discuss the link between the radio properties of the galaxies considered with their stellar mass, with their optical brightness profile and with the presence of a stellar nucleus; we then perform an analysis of the multiwavelength properties of their nuclear emission, taking advantage of X-ray data taken from the literature. Summary and conclusions are given in Section \\ref{summary}. ",
        "conclusions": "\\label{summary}  We presented the results of a radio survey of early-type galaxies in the Virgo cluster, extracted from the sample selected by \\citet{cote04} for their ACS Virgo Cluster Survey. We observed 56 galaxies at 8.4 GHz with the VLA that, combined with data from the literature, provide radio images for all 63 galaxies brighter than B$_{\\rm T} = 14.4$. The aim of this survey is to investigate the origin of radio emission in early-type galaxies and its link with the host properties in an unexplored territory toward the lowest levels of both radio and optical luminosities.  Compact radio sources are found in 12 objects, with fluxes from 0.13 to 2700 mJy. The remaining 51 galaxies are undetected at a flux limit of $\\sim$0.1 mJy, corresponding to radio luminosities L $\\lesssim 4 \\times 10^{18} \\, {\\rm W Hz}^{-1}$. The fraction of radio-detected galaxies is a strong function of stellar mass, in agreement with previous results on the bivariate radio/optical luminosity function of early-type galaxies: while none of the 30 galaxies with M$_{\\star} < 1.7 \\times 10^{10} M_{\\sun}$ is detected, 8 of the 11 most massive galaxies have radio cores.  However, the galaxy mass is not a good predictor of its level of radio emission. This is clearly seen from a comparison of the radio properties of the brightest galaxies of the VCC. While they span only a factor of 10 in stellar mass, they cover a range of $\\gtrsim$ 4.5 orders of magnitude in radio-core power, and three of them are not detected in our radio images.  We note that VCC early-type galaxies with strong dust lanes or of signs of interactions were excluded by the original definition of the sample for the ACS Virgo cluster survey. Considering the links between mergers, dust content and nuclear activity suggested by previous studies of early-type galaxies \\citep[e.g.,][]{tran01,deruiter02,lauer05} the sub-sample considered might be biased against active galaxies with respect to the overall population. This effect, however, does not affect the bright luminosity end, since the sample is complete down to B$_{\\rm T}$ = 12.  We examined the possibility of a link between nucleation and radio emission. Core galaxies lack stellar nuclei and show a high incidence of radio nuclei. In the rest of the sample, only two nucleated galaxies show the presence of a radio source. Thus nuclear activity can coexist with a stellar nucleus, but the fraction of active galaxies is not related to nucleation.  We considered the properties of the surface brightness profiles, separating galaxies reproduced by a S\\'ersic law from those better described by a core-S\\'ersic profile. Furthermore, we relied on the results of X-ray {\\it Chandra} observations by \\citet{gallo08} to derive a multiwavelength view of the AGN emission. The reported link between optical surface brightness profile and radio loudness \\citep{paper3} is found also in the VCC sample, since RL AGNs are associated with core galaxies, while non-core galaxies only host RQ AGNs.  Not all radio-detected galaxies have a X-ray nuclear counter part, and vice-versa. Not surprisingly, it is easier to detect the AGN emission in the radio band for a RL source and in the X-ray band for a RQ AGNs. This has an important consequence when we seek for a complete census of AGNs and consequently of super-massive black holes for which a combination of (at least) radio and X-ray data is required.  Nonetheless, nuclear emission is not detected in a significant fraction of VCC galaxies in either observing band. Clearly, there is the possibility that the VCC sample crosses the minimum galaxy mass at which a SMBH is present (if such threshold indeed exists). However, there are no signs for the presence of an AGN in relatively massive VCC galaxies (e.g., VCC~798 and VCC~1720) despite the fact that they are expected to host a large SMBH (M$_{\\rm BH} \\sim 10^8M_{\\sun}$), based on a typical ratio of 0.002 between SMBH and galaxy's mass \\citep{marconi03}. This is particularly worrisome for a general use of AGNs as black-hole tracers, considering the proximity of the Virgo cluster, the depth and the spatial resolution of the available observations.  In general, we cannot distinguish for the quiescent galaxies between a scenario 1) of a transition toward galaxies lacking a SMBH due to a low mass threshold, 2) the possibility that the SMBH has been ejected due to the recoil caused by the coalescence between two black holes following a galaxies merger \\citep[e.g.,][]{campanelli07}, and 3) of extremely low accretion rate.  In fact, a crucial role in our ability to detect an active nucleus, and in setting its level of activity, is certainly played by the level of accretion onto the central SMBH. At least for low-power RL AGNs, it has shown that a suggestive linear correlation exists between the accretion rate of hot gas (estimated in a spherical approximation) and the jet power, over several orders of magnitude \\citep{allen06,balmaverde08}. Therefore, to take full advantage of active nuclei as black hole tracers, it is necessary to combine a multiwavelength approach with estimates of the accretion rate. It is certainly of great interest to explore for example whether the connection between the level of accretion and activity can be extended to the lowest level of radio emission seen in the RL VCC galaxies, and how this compares with similar estimates for RQ AGNs."
    },
    "0910/0910.3208_arXiv.txt": {
        "abstract": "We present a new observational method to type the explosions of young supernova remnants (SNRs). By measuring the morphology of the {\\it Chandra} X-ray line emission in seventeen Galactic and Large Magellanic Cloud SNRs with a multipole expansion analysis (using power ratios), we find that the core-collapse SNRs are statistically more asymmetric than the Type Ia SNRs. We show that the two classes of supernovae can be separated naturally using this technique because X-ray line morphologies reflect the distinct explosion mechanisms and structure of the circumstellar material. These findings are consistent with recent spectropolarimetry results showing that core-collapse SNe are intrinsically more asymmetric. ",
        "introduction": "Over the past few decades, observations and theory of supernovae explosions (SNe) and their remnants have advanced tremendously. The original classification of SNe was based on their optical spectra near maximum brightness (Minkowski 1941; see Filippenko 1997 for a review), when features originate from the outermost layers of the exploding stars. At later times, SN expansion causes the chemical and physical structure of the SN ejecta and the circumstellar material (CSM) to dominate the spectra. Study of these nebular spectra revealed two physically disinct classes of SNe: Type Ia and core-collapse (CC; Types Ib, Ic, and II) explosions. There is a broad consensus now that Type Ia SNe constitute the thermonuclear detonation of a white dwarf, while CC SNe follow the gravitational collapse of a massive ($>8 M_{\\sun}$) star. These different pathways produce distinct observational signatures that reflect their underlying explosion mechanisms.  As the ejecta continue to expand following an explosion, they are heated to X-ray emitting temperatures by collisionless shocks. Due to this phenomenon, X-ray emission lines are observable in young supernova remnants (SNRs with ages $<$10,000 years) from the metals produced within the stellar interiors and during the SNe. Thus, study of X-ray spectra and X-ray morphology can reveal much about SNe progenitors (see Weisskopf \\& Hughes 2006 for a review). For example, the strength of X-ray emission lines can be used to infer chemical abundances and to compare with the values predicted by Type Ia and CC SNe models. In particular, the O/Fe ratio is often used for this purpose, since Type Ia SNe produce comparatively more Fe than CC SNe, and CC SNe yield large amounts of O. However, the analysis of ejecta-dominated X-ray spectra is very challenging given our limited understanding of a number of topics, including the hydrodynamic interaction between the ejecta and the ambient medium, the inner structure of the ejecta, and the electron heating processes at the reverse shock (Badenes et al. 2003, Rakowski et al. 2006). So far, this detailed analysis of the ejecta emission including all these subtleties has only been done for a few objects.  In this Letter, we present a new, observational method for the typing of young SNRs. Specifically, we describe how the asymmetry of \\hbox{X-ray} line emission from SNRs can be used to discern between a Type Ia or a CC origin. Our technique is general and quantitative, and it is independent of the conditions of the plasma. Thus, this approach can be used systematically for comparison between many sources.  \\begin{deluxetable*}{cccccccc}[ht!] \\tablecolumns{8} \\tablewidth{0pc} \\tabletypesize{\\footnotesize} \\setlength{\\tabcolsep}{0.0in} \\renewcommand{\\arraystretch}{1.0} \\tablewidth{0pt} \\tablecaption{Archival SNRs, Sorted by Age} \\tablehead{\\colhead{Source} & \\colhead{ObsID} & \\colhead{ACIS Exp.} & \\colhead{Age} & \\colhead{Distance} & \\colhead{Radius\\tablenotemark{a}} & \\colhead{$L_{X}$\\tablenotemark{b}} & \\colhead{References} \\\\ \\colhead{} & \\colhead{} & \\colhead{(ks)} & \\colhead{(years)} & \\colhead{(kpc)} & \\colhead{(pc)} & \\colhead{($\\times 10^{37}$ erg/s)} & \\colhead{}} \\startdata \\cutinhead{Type Ia Sources} 0509$-$67.5 & 776, 7635, 8554 & 113 & 350--450 & 50 & 5.96 & 1.76 & \\cite{bad08} \\\\ Kepler & 6714--6718, 7366 & 751 & 405 & 5.0 & 3.88 & 0.19 & \\cite{reynolds} \\\\ Tycho & 3887 &  150 & 437 & 2.4 & 3.72 & 0.12 & \\cite{warren05} \\\\ 0519$-$69.0 & 118 & 40 & 400--800 & 50 & 6.56 & 1.06 & \\cite{rest05} \\\\ N103B & 125 & 37 & 860 & 50 & 5.96 & 1.67 & \\cite{lewis03} \\\\ DEM L71 & 775, 3876, 4440 & 148 & $\\sim$4360 & 50 & 11.9 & 0.77 & \\cite{hughes03} \\\\ 0548$-$70.4 & 1992 & 60 & $\\sim$7100 & 50 & 17.9 & 0.96 & \\cite{hendrick} \\\\ \\cutinhead{Core-collapse Sources} Cas A & 4634--4639, 5196, 5319--5320 & 993 & 309--347 & 3.4 & 3.77 & 2.58 & \\cite{hwang04} \\\\ W49B & 117 & 55 & $\\sim$1000 & 8.0 & 6.30 & 4.48 & \\cite{lal} \\\\ G15.9$+$0.2 & 5530, 6288, 6289 & 30 & $\\sim$1000 & 8.5 & 8.72 & 1.38 & \\cite{reynolds06} \\\\ G11.2$-$0.3 & 780--781, 2322, 3909-3912 & 95 & 1623 & 5.0 & 4.17 & 1.18 & \\cite{kaspi01} \\\\ Kes 73 & 729 & 30 & 500--2200 & 8.0 & 6.11 & 0.94 & \\cite{gotthelf} \\\\ RCW 103 & 970 & 49 & $\\sim$2000 & 3.3 & 5.43 & 2.21 & \\cite{carter} \\\\ N132D & 5532, 7259, 7266 & 90 & $\\sim$3150 & 50 & 21.47 & 9.92 & \\cite{bork07} \\\\ G292.0$+$1.8 & 6677--6680, 8221, 8447 & 516 & $\\sim$3300 & 6.0 & 9.45 & 0.56 & \\cite{park07} \\\\ N49B & 1041 & 35 & 10000 & 50 & 23.85 & 2.39 & \\cite{park03} \\\\ B0453$-$685 & 1990 & 40 & 13000 & 50 & 20.27 & 0.54 & \\cite{gaensler03} \\\\ \\enddata \\tablenotetext{a}{Radius $R$ used in Eq. 3 of the power-ratio calculation. This radius is selected to enclose the entire source in the full-band X-ray image, and it is determined assuming the distances listed above.} \\tablenotetext{b}{X-ray luminosity in the 0.5--2.1 keV band, from the {\\it Chandra} SNR Catalog. The value for G15.9$+$0.2 is calculated using an absorbed planar shock (vpshock) model of the integrated spectra \\citep{reynolds06}.} \\end{deluxetable*} ",
        "conclusions": "The calculated power ratios $P_2/P_0$ and $P_3/P_0$ for the Si {\\sc xiii} images of our sources are plotted in Figure~\\ref{fig:p2p3} (left), with example Si {\\sc xiii} images of seven sources shown in Figure~\\ref{fig:p2p3} (right). The SNRs with elongated or barrel-like features (like W49B, Cas A, and G292.0$+$1.8) have the largest $P_{2}/P_{0}$ values, and sources that are more compact and circular (e.g., 0519$-$69.0 and DEM L71) have the lowest $P_{2}/P_{0}$. Some sources that appear spherical by eye (like 0509$-$67.5 or 0548$-$70.4) have large $P_{2}/P_{0}$ because they are brighter on one side and their full-band centroids have off-center positions. SNRs with obvious mirror asymmetries (such as N132D and N103B) produce large $P_{3}/P_{0}$, while sources with line emission distributed evenly around their full-band centroids (e.g., DEM L71 and 0509$-$67.5) have small \\hbox{$P_{3}/P_{0}$}.  Overall, we find that the power ratios of the Type Ia SNRs are significantly different than those of the CC SNRs. The CC SNe have roughly an order of magnitude larger quadrupole power ratio: the mean $P_{2}/P_{0}$ of the Type Ia SNe is (7.7$\\pm$0.7)$\\times 10^{-7}$ with a standard deviation of 6.9$\\times10^{-7}$ (excluding SNR 0548$-$70.4, see discussion below) and of the CC SNe is (80.3$\\pm$2.6)$\\times 10^{-7}$ with a standard deviation of 56.1$\\times10^{-7}$. This result implies that the CC SNR line emission is comparatively much more elliptical or have more off-center centroids than Type Ia SNRs. We find that the mean $P_{3}/P_{0}$ are also different: the mean of the Type Ia SNe is (1.7$\\pm$0.1)$\\times 10^{-7}$ (excluding SNR 0548$-$70.4) with a standard deviation of 2.1$\\times10^{-7}$ and of the CC SNe is (3.5$\\pm$0.3)$\\times 10^{-7}$ with a standard deviation of 1.8$\\times10^{-7}$. This result indicates that the X-ray line emission in CC SNe is more mirror asymmetric than Type Ia SNe. Additionally, we find that the Type Ia SNRs with greater ellipticity or off-center centroids (larger $P_{2}/P_{0}$) tend to have larger mirror symmetry (smaller $P_{3}/P_{0}$), and those with less mirror symmetry are significantly more circular or balanced surface brightness (smaller $P_{2}/P_{0}$). Overall, the Type Ia and CC SNe seem to naturally form two distinct populations in the $P_2/P_0$ -- $P_3/P_0$ plane.  Only one source of the seventeen is an outlier in Figure \\ref{fig:p2p3}, SNR 0548$-$70.4. This source has two bright limbs as well as a large, luminous area that is off-center and possibly seen in projection (see the right panel of Figure \\ref{fig:p2p3}). \\cite{hendrick} categorized this source as a Type Ia based on its low O/Fe ratio of $\\sim$16, an intermediate value between the O/Fe ratio of $\\sim$0.75 expected for a Type Ia \\citep{iwa} and the O/Fe ratio of 70 predicted for a 20 $M_{\\sun}$ CC SNe; these authors present several possible explanations for an elevated oxygen abundance in a Type Ia SN. In our opinion, further study of SNR 0548$-$70.4 is warranted to explore its explosion type and anomalous ejecta properties.  By including targets from the LMC and the Milky Way, we aimed to limit any biases in our sample. Our SNe span a large range of ages and X-ray luminosities (see Table 1). On average, our Type Ia sources have younger estimated ages than the CC SNe, but we find no trends in the power ratios with age or with radius ($R$ from Table 1). The 0.5--2.1 keV X-ray luminosities of the Type Ia and CC SNe are approximately similar, so our analysis is not biased by the brightness of the sources. We note here that we have performed the same analyses for other X-ray emission lines (e.g., Ne {\\sc ix}, Mg {\\sc xi}, S {\\sc xv}), and all produce the the two distinct populations in the $P_2/P_0$ -- $P_3/P_0$ plane.  Our results show that as a class, CC SNRs are much more asymmetric or elliptical than Type Ia SNRs. These morphological differences reflect the distinct explosions and CSM structures of the progenitors of Type Ia and CC SNe. This finding is consistent with the emerging picture from spectropolarimetry studies that CC SNe are intrinsically aspherical (see Wang \\& Wheeler 2008 for a review), whereas Type Ia SNe are largely spherical (with some variation; see Kasen et al. 2009). Further observations and analyses are necessary to elucidate whether the asymmetries in CC SNRs are properties determined by the central engine/ignition process.  As one possible end of the stellar life cycle, SNe provide crucial tests of stellar evolution models. However, in order to compare observation with theory, we need viable progenitor diagnostics for each SN type. During and immediately after explosions, many aspects of the physical mechanism can be explored: e.g., neutrinos, polarization, radioactive decays, relic compact objects. As presented here, the distribution of line emission in young SNRs gives a complementary and unique means to probe the explosion mechanism many years after a SN event. This tantalizing possibility has always been the magnet of young SNR studies, and the superb spatial and spectral resolution of {\\it Chandra} has enabled detailed analyses of the metal-rich ejecta in SNRs. The clean, quantitative method we introduce here is based entirely on observables, and it represents the first standardized approach to type young SNRs.  \\vspace{5 mm}  We thank the referee, Patrick Slane, for his helpful feedback. This work is supported by DOE SciDAC DE-FC02-01ER41176 (LAL and ER-R) and a National Science Foundation Graduate Research Fellowship (LAL)."
    },
    "0910/0910.3949_arXiv.txt": {
        "abstract": "We try to address quantitatively the question whether a new mass is needed to fit current supernovae data. For this purpose, we consider an infra-red modification of gravity that does not contain any new mass scale but systematic subleading corrections proportional to the curvature. The modifications are of the same type as the one recently derived by enforcing the ``Ultra Strong Equivalence Principle\" (USEP) upon a Friedmann-Lema\\^{i}tre-Robertson-Walker universe in the presence of a scalar field. The distance between two comoving observers is altered by these corrections and the observations at high redshift affected at any time during the cosmic evolution. While the specific values of the parameters predicted by USEP are ruled out, there are regions of parameter space that fit SnIa data very well. This allows an interesting possibility to explain the apparent cosmic acceleration today without introducing either a dark energy component or a new mass scale. ",
        "introduction": "During the last decade, several observational probes \\cite{SNIa,CMB,BAO} have confirmed that our universe is undergoing a phase of accelerated expansion. Beyond the details of specific models, one of the most remarkable aspects of such a discovery is the seemingly unavoidable presence of a new tiny mass scale in the theory that describes our world.  In the framework of General Relativity (GR), a negative pressure component (``dark energy''~\\cite{review}) can account for the cosmic acceleration. During the cosmological expansion, such a component has to become dominant when the average energy density $\\rho(t)$ drops to about its present value $\\rho_0$ ($t$ is the proper time). Thus, dark energy Lagrangians typically contain a mass parameter of the order of $M \\sim \\rho_0^{1/4} \\sim 10^{-3} {\\rm eV}$, that triggers the epoch when dark energy starts to dominate\\footnote{The mass parameter can be higher e.g. in quintessence models with power-law potentials~\\cite{Steinhardt:1999nw,Ratra}, but at the price of giving a milder equation of state which is now severely challenged by observations.}. Models of massive/modified gravity highlight the problem from a different perspective. If the graviton is effectively massive, the modified dynamics of gravity at large distances can provide a mechanism for ``self-acceleration\" \\cite{dgp} and/or of filtering for the cosmological constant's zero mode~\\cite{justin}. In that case the new mass brought into the theory is the mass of the graviton\\footnote{Interestingly, and as opposed to, e.g., the mass of scalar quintessence fields, such a mass might be protected against -- and actually made smaller by -- radiative corrections~\\cite{dvali}.}, $m_g$, typically of order the Hubble constant $H_0$, i.e. $m_g \\sim H_0 \\sim (\\rho_0/M_{\\rm Pl}^2)^{1/2} \\sim 10^{-33} {\\rm eV}$. In $f(R)$ theories of gravity, where the Lagrangian density $f$ is a function of the Ricci scalar $R$, the late-time acceleration is realized, again, by introducing a curvature scale $R_c$ of the order of $H_0^2$~\\cite{fRviable} (see also Refs.~\\cite{fRviable2}). Note that, although based on different mechanisms, all models follow an analogous pattern\\footnote{The few alternatives to this common pattern include the proposal that our universe is not homogeneous on large scales~\\cite{inho} and attempts based on possible non-trivial effects of smaller inhomogeneities on the cosmic evolution~\\cite{back}.}: there is a tuned scale ``hidden\" in the theory which becomes effective, ``by coincidence\", when the appropriate cosmological quantity (the average density $\\rho(t)$ or the Hubble parameter $H(t)$) drops to about its value.  But is the acceleration of the universe intrinsically implying the presence of a new mass scale? If we allow the possibility of departures from GR at large distances there is a logical alternative. High redshift observations have the unique property of relating objects (e.g., the observer and the supernova) that are placed from each other at a relative distance of the order of the average inverse curvature (roughly, the Hubble length $\\sim H_0^{-1}$). Therefore, modifying GR in the infrared (IR) at a length scale set by the curvature -- rather than fixed \\emph{a priori} by a parameter -- will systematically affect any cosmological observation at high redshift, regardless of when such an observation takes place and without the need of any external mass scale. In other words, we might not need a new mass scale because we already have (a dynamical) one, the Hubble parameter $H(t)$; the only ``coincidence\" that we might be experiencing is that of observing objects that are placed from us as far as the Hubble radius\\footnote{The same circumstance does not apply, for instance, to observations within the solar system: typical solar system distances are always extremely small with respect to, say, the average Weyl curvature. An order of magnitude estimate indeed gives $({\\rm Weyl\\ curvature})^{-1/2} \\approx r\\,(r/3\\,{\\rm km})^{1/2}$, where $r$ is the distance from the Sun.}.  The point of view sketched above is somewhat compelling, it addresses directly the fine tuning and coincidence problems, but seems to require a serious revision of the current low energy framework for gravity. Any gravitational operator that becomes effective in the infrared, on dimensional grounds, has to bring in the Lagrangian a mass parameter. Moreover, GR itself is already a geometrical deformation of flat space at distances of the order of the curvature. What seems to be required is a further curvature-dependent subleading effect that systematically modifies the geometrical description of GR at large distances.  Recently, a proposal along those lines has been made by one of the present authors. The modification upon the standard framework is forced by imposing an ``Ultra Strong\" version of the Equivalence Principle (USEP, see Refs.~\\cite{Piazza1,fedo} for more details). USEP suggests that the usual geometric description of spacetime as a metric Riemannian manifold might hold only approximately, at small distances. Such a conjectured ``IR-completion\" of gravity, in its full generality, represents a major theoretical challenge. However, it can be tentatively explored with a Taylor expansion around GR, by applying USEP to a specific GR solution\\footnote{In a similar fashion, someone who does not know GR can try to expand around some point in Riemann normal coordinates and find, in some specific cases, the first GR corrections to flat space.} (see Appendix \\ref{sketch}). For a spectator scalar field in a spatially flat Friedmann-Lema\\^{i}tre-Robertson-Walker (FLRW) universe, the first-order correction to GR is calculable, and few cosmological consequences are derivable~\\cite{fedo}.  Consider, as the  zeroth-order (GR-) approximation, a homogeneous, spatially flat FLRW universe with scale factor $a(t)$. It is known that in such a solution the physical distance $d(t)$ between a pair of comoving observers grows proportionally to $a(t)$; otherwise stated, the comoving distance $\\lambda \\equiv d(t)/a(t)$ is a constant. The first-order correction found in Ref.~\\cite{fedo} modifies such an expansion law by a subleading distance-dependent contribution. As a result, the distance $d(t)$ between two comoving observers grows as $a(t)$ only when  small compared to the Hubble length $H^{-1}$ but gets relevant corrections otherwise. Thus, the scale factor $a(t)$ defines the expansion everywhere but only in the local limit, and effectively detaches from the expansion on the largest scales.  The corrected expansion is most easily seen in terms of the above defined comoving distance $\\lambda$, which is constant only in the small distance limit. Its derivative with respect to observers' proper time $t$ reads~\\cite{fedo} \\begin{equation} \\dot{\\lambda}=\\lambda^3 \\frac{(H^2 a^2)^{\\cdot}}{4}\\,+ {\\rm higher\\ orders}, \\label{dotlamre} \\end{equation} and is clearly negligible on sub-Hubble scales. For completeness, a basic derivation of Eq.~(\\ref{dotlamre}) is sketched in Appendix~\\ref{sketch}. The comoving trajectory $r(t)$ of a light ray also receives corrections because the modified global expansion (\\ref{dotlamre}) has to be considered on top of the usual contribution $\\rd r = \\rd t /a$. In a matter-dominated universe, where the Hubble parameter at the redshift $z=1/a-1$ is given by $H(z)=H_0(1+z)^{3/2}$, we have \\begin{equation} \\frac{\\rd (H_0 r)}{\\rd z}=\\frac{1}{(1+z)^{3/2}} +\\frac{(H_0 r)^3}{4}\\,, \\label{redshift} \\end{equation} which has to be solved with initial conditions $r(0) = 0$.  The above modification, including the factor of $1/4$, is forced by requiring that USEP applies for a scalar field in a spatially flat FLRW universe~\\cite{fedo}. The correction increases the luminosity distance \\begin{equation} d_L(z)=(1+z)r(z)\\,, \\label{luminosity} \\end{equation} and therefore it effectively goes in the direction of a universe with positive acceleration. However, as the present analysis shows as a by-product, the correction given in the last term of Eq.~(\\ref{redshift}) is too small (of too high order in the redshift) to explain SnIa data.  Equation (\\ref{redshift}) is an example of a IR-geometrical deformation that does not contain any mass scale but only $H$-dependent subleading terms. In this paper, by studying a generalized version of (\\ref{redshift}), we attempt to address quantitatively the question whether SnIa data can be explained without introducing any new mass scale. We include terms of lower order in the redshift that are needed to efficiently reproduce SnIa observations.  Equation (\\ref{dotlamre}) is generalized as follows: \\begin{eqnarray} \\dot \\lambda & = & A_1\\,  \\lambda H + A_2 \\,  \\lambda^2 H^2 a + \\dots \\nonumber \\\\ & & +B_1 \\, \\lambda^2 (H a)^{\\cdot} + B_2 \\, \\lambda^3  (H^2 a^2)^{\\cdot} + \\dots\\, . \\label{dotlam} \\end{eqnarray} Note that the above structure of corrections includes (\\ref{dotlamre}) as a special case. In practice, the terms in the above expansion rearrange when we calculate the luminosity distance. Therefore, for a matter-dominated universe, a quite general structure of subleading terms is given by \\begin{equation} \\label{ODE} \\frac{\\rd (H_0 r)}{\\rd z} = \\frac{1}{(1+z)^{3/2}} F\\left(H_0 r (1+z)^{1/2}\\right), \\end{equation} where $F(x)$ is a generic function with $F(0) = 1$: \\begin{equation} \\label{ODE2} F(x) = 1 + \\alpha x + \\beta x^2 + \\gamma x^3+\\, \\dots \\, . \\end{equation} By comparison with (\\ref{dotlam}) we have $ \\alpha=-A_1, \\beta=B_1/2-A_2, \\gamma=B_2\\,.$ Note that Eq.~(\\ref{redshift}) corresponds to $\\alpha = \\beta = 0$ and $\\gamma=1/4$, while in GR all coefficients are set to zero. ",
        "conclusions": "\\label{conclusions}      We have considered IR modifications of gravity that do not imply the presence of a new mass scale in the theory and we have studied their compatibility with the SnIa data. Our first result is that the mechanism derived in Ref.~\\cite{fedo} (see also Appendix \\ref{sketch}), Eq.~(\\ref{dotlamre}), is not enough, by itself, to describe the observed amount of acceleration. The absence of free parameters in equation (\\ref{dotlamre}) (the model in~\\cite{fedo} has one parameter less than $\\Lambda$CDM) does not make up for the very poor fit of the data. However, a more general structure of corrections (\\ref{dotlam}) can lead to a sensibly larger luminosity distance than in the Einstein de Sitter universe. In particular, when $\\gamma=0$ in Eq.~(\\ref{ODE})-(\\ref{ODE2}), we have found that the model fits the data very well for the values close to the exact numbers $\\alpha=1$ and $\\beta=-1/2$. It is interesting to consider such sharp numerical values, not because of abstract numerology, but because a mechanism analogous to that described in Appendix~\\ref{sketch} very naturally produces coefficients which are integers or simple fractions.  At present it is not clear if the corrections that one finds by applying USEP to a field  automatically apply to, or are inherited by, other types of fields. The suggested luminosity distance may eventually turn out to be produced by considering other fields\\footnote{It is interesting, for instance, that a different mechanism, based on a Casimir-like vacuum energy~\\cite{urban}, needs a Veneziano ghost in order reproduce a density of  the right order of magnitude.} than the scalar field considered in~\\cite{fedo}  or by means of other theoretical insights. We also considered the full three parameters model (\\ref{ODE})-(\\ref{ODE2}), whose best fit considerably improves the $\\chi^2$ and   is found to be favored over $\\Lambda$CDM by the AIC but not by the BIC criterion.  It should be noted that the parameters that we are fitting are not coupling constants, and do not appear in a Lagrangian. Rather, they are intended as the terms of a series expansion that approximate the new ``IR-completed\" theory starting from GR. As mentioned in the introduction, the proposed deformation is present at any time during the cosmological evolution; it affects any cosmological observation at high redshift, regardless of when such an observation takes place, and therefore addresses the ``coincidence problem\" in the most direct way.     We note that our model requires a rather low value of the Hubble constant, $H_0 \\sim 50$\\,km/s/Mpc, compared to the constraint from the Hubble Key Project with the determination of $H_0=72 \\pm 8$\\,km/s/Mpc \\cite{Freedman}. This comes from the fact that the Hubble parameter evolves as $H(z) \\simeq H_0 (1+z)^{3/2}$ even for $z \\lesssim {\\cal O}(1)$ due to the absence of a dark energy component. However, methods of the determination of $H_0$ that are largely independent of distance scales of the Large Magellanic Cloud Cepheid typically give low values of $H_0$ \\cite{Sarkar}. For example, Reese {\\it et al.} \\cite{Reese} showed that Sunyaev-Zel'dovich distances to 41 clusters provide the constraint $H_0=54^{+4}_{-3}$\\,km/s/Mpc in the Einstein de Sitter universe. Thus the information of $H_0$ alone is not yet sufficient to rule out our model.  It would be interesting to see the effect of the proposed IR correction on the angular diameter distance to the last scattering surface and estimate the modification to the temperature anisotropies in Cosmic Microwave Background (CMB). CMB data will be certainly useful to place further constraints on our model and to complete the picture at higher redshift."
    },
    "0910/0910.4875_arXiv.txt": {
        "abstract": "We present high-precision photometry of five consecutive transits of WASP-18, an extrasolar planetary system with one of the shortest orbital periods known. Through the use of telescope defocussing we achieve a photometric precision of 0.47--0.83 mmag per observation over complete transit events. The data are analysed using the {\\sc jktebop} code and three different sets of stellar evolutionary models. We find the mass and radius of the planet to be $M_{\\rm b} = 10.43 \\pm 0.30 \\pm 0.24$\\Mjup\\ and $R_{\\rm b} = 1.165 \\pm 0.055 \\pm 0.014$\\Rjup\\ (statistical and systematic errors) respectively. The systematic errors in the orbital separation and the stellar and planetary masses, arising from the use of theoretical predictions, are of a similar size to the statistical errors and set a limit on our understanding of the WASP-18 system. We point out that seven of the nine known massive transiting planets ($M_{\\rm b} > 3$\\Mjup) have eccentric orbits, whereas significant orbital eccentricity has been detected for only four of the 46 less massive planets. This may indicate that there are two different populations of transiting planets, but could also be explained by observational biases. Further radial velocity observations of low-mass planets will make it possible to choose between these two scenarios. ",
        "introduction": "The recent discovery of the transiting extrasolar planetary system WASP-18 \\citep[][hereafter H09]{Hellier+09natur} lights the way towards understanding the tidal interactions between giant planets and their parent stars. WASP-18\\,b is one of the shortest-period ($P_{\\rm orb} = 0.94$\\,d) and most massive ($M_{\\rm b} = 10$\\Mjup) extrasolar planets known. These properties make it an unparallelled indicator of the tidal dissipation parameters \\citep{GoldreichSoter66icar} for the star and the planet \\citep{Jackson+09apj}. A value similar to that observed for Solar system bodies ($Q \\sim 10^4$--$10^6$; \\citealt{Peale99araa}) would cause the orbital period of WASP-18 to decrease at a sufficient rate for the effect to be observable within ten years (H09).  In this work we present high-precision follow-up photometry of WASP-18, obtained using telescope-defocussing techniques \\citep{Me+09mn,Me+09mn2} which give a scatter of only 0.47 to 0.83 mmag per observation. These are analysed to yield improved physical properties of the WASP-18 system, with careful attention paid to statistical and systematic errors. The quality of the light curve is a critical factor in measurement of the physical properties of transiting planets \\citep{Me09mn}. In the case of WASP-18 the systematic errors arising from the use of theoretical stellar models are also important, and are a limiting factor in the understanding of this system. ",
        "conclusions": "We have presented high-quality observations of five consecutive transits by the newly-discovered planet WASP-18\\,b, which has \\reff{one of the shortest orbital periods} of all known transiting extrasolar planetary systems (TEPs). Our defocussed-photometry approach yielded scatters of between 0.47 and 0.83 mmag per point in the final light curves. These data were analysed using the {\\sc jktebop} code, which was modified to include the spectroscopically derived orbital eccentricity in a statistically correct way. The light curve parameters were then combined with the predictions of theoretical stellar evolutionary models to determine the physical properties of the planet and its host star.  A significant source of uncertainty in our results stems from the use of theoretical models to constrain the physical properties of the star. Further uncertainty comes from observed \\Teff\\ and \\MoH, for which improved values are warranted. However, the systematic error from the use of stellar theory is an important uncertainty in the masses of the star and planet. This is due to our fundamentally incomplete understanding of the structure and evolution of low-mass stars. As with many other transiting systems (e.g.\\ WASP-4; \\citealt{Me+09mn2}), our understanding of the planet is limited by our lack of understanding of the parent star.  \\reff{We confirm and refine the physical properties of WASP\\,18 found by H09.} WASP-18\\,b is a very massive planet in an extremely short-period and {\\em eccentric} orbit, which is a clear indicator that the tidal effects in planetary systems are weaker than expected (see H09). Long-term follow-up studies of WASP-18 will add progressively stricter constraints on the orbital decay of the planet and thus the strength of these tidal effects.  We now split the full sample of known (i.e.\\ published) TEPS into two classes according to planetary mass. The mass distribution of transiting planets shows a dearth of objects with masses in the interval 2.0--3.1\\Mjup. There are nine planets more massive than this and 46 less massive. Seven of the nine high-mass TEPs have eccentric orbits (HAT-P-2, \\citealt{Bakos+07apj3}; HD\\,17156, \\citealt{Barbieri+07aa}; HD\\,80606, \\citealt{Laughlin+09natur}; WASP-10, \\citealt{Christian+09mn}; WASP-14, \\citealt{Joshi+09mn}; WASP-18; XO-3, \\citealt{Johnskrull+08apj}), and the existing radial velocity observations of the remaining two cannot rule out an eccentricity of $0.03$ or lower (CoRoT-Exo-2, \\citealt{Alonso+08aa2}; OGLE-TR-L9, \\citealt{Snellen+09aa}). By comparison, only four of the 46 low-mass TEPs have a {\\em significant} \\citep{LucySweeney71aj} orbital eccentricity measurement.  These numbers imply that the more massive TEPs are a different population to the less massive ones; Fisher's exact test \\citep{Fisher22jrss} returns a probability lower than $10^{-5}$ of the null hypothesis (although this does not account for our freedom to choose the dividing line between the two classes). This indicates that the two types of TEPs have a different \\reff{internal structure, formation mechanism, or evolution,} a suggestion which is supported by observations of misalignment between the spin and orbital axes of $M > 3$\\Mjup\\ TEPs \\citep{Johnson+09pasp}.  There is, however, a bias at work here. The more massive TEPs cause a larger radial velocity signal in their parent star ($M_{\\rm b} \\propto K^3$), so a given set of radial velocity measurements can detect smaller eccentricities \\citep[see also][]{ShenTurner08apj}. The eccentricity of the WASP-18 system is in fact below the detection limit of existing observations of most TEPs. We therefore advocate the acquisition of additional velocity data for the known low-mass TEPs, in order to equalise the eccentricity detection limits between the two classes of TEPs. These observations would allow acceptance or rejection of the hypothesis that more massive TEPs represent a fundamentally different planet population to their lower-mass brethren."
    },
    "0910/0910.1985_arXiv.txt": {
        "abstract": "{}{We analyze observational data from 4 instruments to study the correlations between chromospheric emission, spanning the heights from the temperature minimum region to the middle chromosphere, and photospheric magnetic field.} {The data consist of radio images at 3.5~mm from the Berkeley--Illinois--Maryland Array (BIMA), UV images at 1600~\\AA\\ from TRACE, \\ca-line filtergrams from BBSO, and MDI/SOHO longitudinal photospheric magnetograms. For the first time interferometric millimeter data with the highest currently available resolution are included in such an analysis. We determine various parameters of the intensity maps and correlate the intensities with each other and with the magnetic field.} {The chromospheric diagnostics studied here show a pronounced similarity in their brightness structures and map out the underlying photospheric magnetic field relatively well. We find a power law to be a good representation of the relationship between photospheric magnetic field and emission from chromospheric diagnostics at all wavelengths. The dependence of chromospheric brightness on magnetic field is found to be different for network and internetwork regions.}{} ",
        "introduction": "The solar chromosphere is known to be highly structured at different spatial scales and the range of chromospheric inhomogeneities is best revealed when observed at different wavelengths from the extreme ultraviolet to millimeter waves \\citep[see, e.g.,][]{2004IAUS..223..195S}. Even far from activity complexes, in the quiet Sun magnetic field is generally believed to be responsible for chromospheric structuring. This is consistent with the strong correlation found between excess emission in the cores of the \\ca\\ and H resonance lines (the primary chromospheric diagnostic) and the presence and strength of quiet-Sun magnetic flux \\citep[e.g.][]{1959ApJ...130..366L,1989ApJ...337..964S}. Consequently, these resonance lines commonly serve as indicators of changes in the chromospheric structure related to global magnetic activity and heating of the outer atmosphere of the Sun and cool stars \\citep{1991SoPh..134...15R,1989ApJ...337..964S}. In addition to the chromospheric emission related to the magnetic field, \\citet{1987A&A...172..111S,1992A&A...258..507S} has suggested that part of the radiative flux from solar and stellar chromospheres, referred to as basal flux, has a nonmagnetic origin and is believed to stem from acoustic heating of the chromosphere \\citep[see, e.g.][]{1991SoPh..134...15R}.  The primary brightness patterns revealed in two-dimensional images obtained in the ionized calcium H and K lines include bright plages, chromospheric network located at the boundaries of supergranular cells, and bright points in the interior of the cells -- internetwork  grains or \\textit{K2V} bright points. Chromospheric network features are known to display a one-to-one spatial correspondence with regions of enhanced photospheric magnetic field \\citep{1975ApJ...200..747S,1982SoPh...80..227S,1998SoPh..179..253N, 2000A&A...363..279S}, while the (magnetic) origin of the internetwork grains is a controversial issue. From the analysis of simultaneous Kitt Peak \\ca-line spectroheliograms and photospheric magnetic area scans, \\citet{2000A&A...363..279S} showed the association of internetwork \\textit{K2V} bright points with magnetic features, confirming the earlier results of \\citet{1982SoPh...80..227S}. However, from analysis of \\ca-line and magnetic field data from the BBSO, \\citet{1998SoPh..179..253N} have found that not all internetwork bright points are associated with magnetic fields. The existence of 2 distinct classes of cell-interior grains, which differ in temporal behaviour and spatial motion and have a different origin, was revealed by observations with the Swedish Solar Telescope reported by \\citet{1992ASPC...26..161B}, whereas \\citet{1999ApJ...517.1013L} found no obvious relation of the internetwork brightenings to the magnetic field, and \\citet{1996A&A...316..196R} concluded that they are a totally nonmagnetic phenomenon.  Quantitative studies of the relationship between magnetic field and the chromosphere (represented by the calcium K and H-lines), in both the active and quiet solar atmosphere, have been carried out by a number of authors \\citep[e.g.][]{1975ApJ...200..747S,1989ApJ...337..964S,1998SoPh..179..253N,1999ApJ...515..812H,2005MmSAI..76.1018O,2007A&A...466.1131R}. In the quiet-Sun chromosphere \\citet{1975ApJ...200..747S} and \\citet{1998SoPh..179..253N} found a linear correlation between the K-line intensities and the absolute values of the magnetic field strength. \\citet{1989ApJ...337..964S} extended this relationship to solar active regions and suggested that the relation between the K-line excess flux density (after the subtraction of the basal flux from the observed \\ca-line flux) and the magnetic field flux density is best described by a power law with an exponent of about 0.6. Considering 4 types of brightness structures, which included active regions, decaying active regions, the enhanced network and quiet (weak) network, \\citet{1999ApJ...515..812H} ascertained that in both active and quiet Sun the relation between the magnetic flux density and calcium residual intensity follows a power law with an exponent of 0.5. \\citet{2005MmSAI..76.1018O} cited a power-law exponent value of 0.66 for the quiet Sun, while \\citet{2007A&A...466.1131R} obtained a lower exponent of 0.2 for locations in a quiet-Sun region, and higher values of 0.4--0.5 for network locations, depending on the magnetic field threshold value.  \\citet{1991A&A...252..203R} found power-law relations between excess flux densities in stellar chromospheric lines with differing temperature responses, with the exponents increasing with increasing difference between their temperatures of formation. Relations between fluxes originating in the same atmospheric regime were found to be close to linear (i.e., having a power-law exponent of unity). For the quiet Sun the high degree of co-spatiality of the brightness features seen in the \\ca\\ line and the \\textit{TRACE} UV channels has been demonstrated by \\citet{1999ASPC..183..383R}.  Almost all these studies have been limited to the visible and UV spectral range. An exception is the work of \\citet{1991ApJ...383..443L}, who studied the cross-correlation statistics of the brightness variations in the submm (850$\\mu$m) and \\ion{Ca}{ii} K-line images. Their results are confined to the temperature minimum region and the heights of the low chromosphere. No detailed studies have compared millimeter emission, which originates from the middle chromosphere, with other chromospheric indicators. From observations at 350, 850 and 1200 $\\mu$m with the James Clerk Maxwell Telescope at a resolution of about 20\\arcsec, \\citet{1995ApJ...453..511L} reported submillimeter chromospheric brightness structures which resemble the supergranular network, although the resolution was marginal for such a study. Initial qualitative results of the comparison of chromospheric structure, identified in millimeter images obtained with BIMA at 3.5 mm with a resolution of about 10\\arcsec, with calcium images from BBSO and line-of-sight photospheric magnetograms from SOHO/MDI were given in \\citet{2006A&A...456..697W}.  In this paper we present an observational study of the relationships between the photospheric magnetic field and emission at different temperatures from chromospheric heights. We study the chromospheric observational data referring to the temperature minimum region (UV continuum at 1600~\\AA),  low chromosphere (\\ca\\ line) and middle chromosphere (continuum at 3.5 mm). This is the first analysis of this type that includes millimeter observational data. Unfortunately, simultaneously recorded observational data on the chromospheric magnetic field suitable for such an analysis are unavailable. For this reason, we focus on the MDI/SOHO longitudinal photospheric magnetograms and study the dependence of the chromospheric intensity at different wavelengths on the photospheric magnetic field strength.  In Sec. 2 and 3 we describe the data and instruments used for the analysis and provide the details of the data reduction. In Sect.~4 the topography of the quiet chromosphere, chromospheric intensity distributions and cross-correlations between the magnetic field and chromospheric emission are presented. In Sect.~5 we summarize and discuss the results. ",
        "conclusions": "We have considered chromospheric disk-center, quiet-Sun emission at different wavelengths and its dependence on photospheric magnetogram signal, including, for the first time, observations at 3.5~mm. Maps in \\ca\\ core intensity, the TRACE 1600~\\AA\\ channel, BIMA 3.5 mm brightness temperature and (unsigned) MDI magnetograms reveal clearly the same pattern, although at different spatial resolutions. At the resolution of 12\\arcsec\\ millimeter brightness resembles quiet-Sun time-averaged emission from other chromospheric heights, and demonstrates a dependence on the underlying magnetic field similar to that of the other chromospheric diagnostics.  The tightest relation between the chromospheric emission and photospheric magnetic field is displayed by the \\ca\\ line. At all wavelengths a power law gives a lower $\\chi^2$ than a linear fit, although for the 1600~\\AA\\ radiation the difference is marginal. For \\ca\\ this result supports the conclusion reached by \\citet{1989ApJ...337..964S}, but based on active region data, and by \\citet{1999ApJ...515..812H} for both active and quiet-Sun magnetic structures. For \\ca\\ \\citet{1975ApJ...200..747S} and \\citet{1998SoPh..179..253N} found a linear dependence to be adequate. However, the linear dependencies were obtained after excluding weak fields (weaker than 25~G and 20~G, respectively) and for a magnetogram signal not exceeding 120~G and 50~G, respectively. In our study we included both weak fields exceeding $1\\sigma$-noise level of the magnetograms (1.5 or 4~G depending on the amount of averaging), and the highest values reached in the quiet-Sun data (about 230~G at MDI's spatial resolution). This suggests that the non-linearity of the relation found in this work is influenced mostly by weak chromospheric emission that reveals a strong dependence on magnetic fields. For instance, if we set a $B$ threshold of 20~G we cannot distinguish between the quality of the linear and power-law fits for all chromospheric emissions analyzed here. Our conclusions agree with those presented by \\citet{2007A&A...466.1131R}, who found from the POLIS measurements at 1\\arcsec\\ resolution a clear non-linear increase of the calcium core flux with magnetic flux for values $B \\leq 100$~G and a slow increase for higher values. They also demonstrated that the power-law exponent depends strongly on the value of the lower threshold.  The difference between the maximum magnetogram signal reached in our data and that in the data of \\citet{1975ApJ...200..747S} and \\citet{1998SoPh..179..253N} may indicate either that their quiet-Sun regions are even quieter than the one we analyzed, or it may reflect differences in spatial resolution. One more difference between our analysis and the others is that we consider temporally averaged data rather than snapshot images. By using time-averaging we reduce significantly the scatter in the relationship and the noise in the magnetograms, which lead to more reliable power-law fits. By employing snapshot images we get more points in the scatter plots but due to higher magnetogram noise level the power-law exponent is even lower (and less reliable) than for the time-averaged images. One should also note the difference in the analysis techniques between this work and the others applied for building the scatter plots. In this work we employed pixel-to-pixel comparison of the images. However, the results of the fits are similar if we consider the binned chromospheric data averaged within a narrow range of magnetogram values. On the whole, from the data analyzed here we are able to confidently determine non-linear behaviour of 3 chromospheric magnetic-intensity relationships.  The exact power law exponents of the relations between the calcium core emission and the photospheric magnetic field differs from that found by \\citet{1989ApJ...337..964S} (0.3 vs. 0.6) . This probably has to do with the smaller $<B>$ range available in the quiet Sun, although it cannot be ruled out that the \\ca\\ intensity does not increase so rapidly in the quiet Sun as in active regions \\citep[see][]{1991A&A...250..220S}. To compare these explanations, we studied the dependence of the calcium emission on the magnetic flux in an active region close to disk center observed by the BBSO and MDI on 11.03.2004. The exponent obtained from a power-law fit for the AR emission is surprisingly low and similar to the QS values, being in the range 0.28-0.32, with the exact value depending somewhat on the criterion chosen to identify the AR pixels on the map. The use of a pixel size and a noise threshold similar to that of \\citet{1989ApJ...337..964S} does not increase the value of the power-law exponent and it still differs from that of \\citet{1989ApJ...337..964S} and \\citet{1999ApJ...515..812H}. The difference may have to do with the spectral range covered by the different observations.  The reduction in the power-law exponent as compared to the value derived by \\citet{1989ApJ...337..964S} may be to a large extent the result of the inclusion of weak magnetic flux regions, which are systematically characterized by magnetogram field strengths that are smaller than the true values due to the presence of weak neutral lines with null line-of-sight component. Unfortunately, we cannot exclude the influence of these horizontal magnetic fields, which are missed by MDI observations. Such fields, partly seen in the internetwork regions by Hinode \\citep[e.g.][]{2007PASJ...59S.571L}, are very likely small loops \\citep[][]{2007ApJ...666L.137C,2007A&A...469L..39M} and may be due to a local dynamo \\citep[][]{2007A&A...465L..43V,2008A&A...481L...5S}. It is unclear if such low-lying loops really influence the structure of the higher--lying chromosphere. However, if they do contribute to the heating of those layers then this could explain the high chromospheric emission from many spatial locations associated with small magnetic fluxes in the MDI data (Fig.~\\ref{fig6}). Unresolved opposite magnetic polarities within the same individual MDI pixels have a similar effect. Both these effects together could explain the result that the power law exponent is lower when very quiet (low flux) pixels are included. Studies at higher spatial resolution and sensitivity are needed to test these hypotheses.  On the whole, chromospheric emission displays a rather different dependence on the network and internetwork magnetic field. The power-law exponents in the relation between the \\ca-line intensity and the network magnetic flux are found to be similar to those obtained by \\citet{1989ApJ...337..964S} for an active region and by \\citet{1999ApJ...515..812H} for the quiet Sun network. The same conclusion holds for the relation between the UV emission at 1600~\\AA\\ and magnetic flux in the network regions of the quiet Sun. For chromospheric fluxes at a resolution of 4\\arcsec\\ it was found that emission from the network entirely defines the trend for magnetic flux densities stronger than $\\simeq 20$~G. Our results for the calcium core emission remarkably confirm those of \\citet{2007A&A...466.1131R}, who obtained for the ensemble of quiet-Sun points a low value for the power-law exponent of $\\simeq$0.2 (for snapshot images), while for the network emission from regions with $B \\geq 20$ G the exponent reaches the value of 0.51 (and 0.45 for $B \\geq 10$ G).   A considerable scatter about the mean relation between the millimeter flux at 3.5~mm and the magnetic flux density per 12\\arcsec\\ aperture makes a confident study of the mean trend difficult. Although the data strongly suggests that the dependence of the 3.5~mm brightness on magnetic field is non-linear and best described by a power law (among the functional forms considered here), but the sizeable scatter does not allow the exponents to be determined reliably. The best fit exponent depends strongly on the range of $<B>$ considered.  In the internetwork regions almost no dependence of the \\ca-line, UV at 1600~\\AA\\ and 3.5~mm intensity on the magnetogram signal was found. Unfortunately, due to the very low $B$ values in the internetwork, only a few percent of IN structures are associated with magnetogram signal exceeding the noise level. However, it should be noted that analysis of the cell interiors and their relation to the underlying magnetic field carried out in \\citet{2007A&A...466.1131R} for \\ion{Ca}{ii} H-line with lower noise and higher resolution has revealed similar results. This supports the idea that the role of magnetic field is negligible for bright internetwork grains (e.g. acoustic heating).  Correlation analysis performed for the quiet-Sun chromospheric emissions surprisingly revealed a linear relation between the \\ca-line intensity and UV brightness at 1600~\\AA. According to \\citet{1991A&A...252..203R}, who claimed that the relations between the fluxes originating in the same atmospheric regimes are commonly close to linear, these chromospheric emissions are both generated in the chromospheric temperature regime from chromospheric ranges of heights. However, solar models place UV emission at 1600~\\AA\\ in the temperature minimum region, while calcium emission is believed to be generated from heights more typical of the lower chromosphere. The power-law relations between calcium flux and 3.5~mm flux, as well as between 3.5~mm brightness and 1600~\\AA\\ flux, suggest different temperatures of formation for these pairs of emissions, but as the power-law exponent do not differ much from unity we may assume that these emissions are still close in the heights of formation. However, the flux correlation analysis for the considered emissions and its conclusions should be treated with caution. Due to the very different temperature sensitivity of, on the one hand, UV and calcium intensity, and, on the other hand, 3.5 mm brightness, they represent different contributions from small (unresolved) hot magnetic features and cool weak-field regions within the same resolution element. In other words, in an inhomogeneous atmosphere the UV samples mainly the hot features, while the mm is composed of radiation from both. This could explain some of the difference in the behaviour of 1600~\\AA\\ and \\ca\\ excess emission compared to mm emission, and, in particular, the large scatter in the flux-flux plots (Fig.~\\ref{fig13}).  Summarizing, clear spatial correlations found between photospheric magnetograms, the \\ca-line, UV intensity at 1600~\\AA\\ and 3.5~mm emission indicate that heating in the quiet-Sun lower and middle chromosphere maps out the underlying photospheric magnetic field rather well. Our results imply that at the chromospheric heights covered by the wavelengths investigated here we deal with the same heating mechanism. However, any mechanism that explains chromospheric heating must explain the relationship between chromospheric intensity excess and magnetic fields. In particular, it should validate both the dominating role of magnetic field in the chromospheric emission of the network and the absence of correlation between the internetwork emission and magnetic flux.  Of particular interest is the possibility of studying the middle chromosphere at the heights where mm emission originates, but because of very limited spatial resolution and insufficient sensitivity a detailed study like that cannot be conducted at present. In general, present data on both chromospheric emission and magnetic field are not sufficient to establish physical mechanisms that are acting at chromospheric heights leading to the observed brightness structures. In this respect, an extraordinarily powerful tool to study the thermal structure of the solar atmosphere from the temperature minimum region to the middle chromosphere will be provided by the Combined Array for Research in Millimeter--wave Astronomy (CARMA) and particularly by Atacama Large Millimeter Array (ALMA) as was demonstrated in \\citet{2008Ap&SS.313..197L}."
    },
    "0910/0910.2462.txt": {
        "abstract": "% context {} % aims { We present a new homogeneous set of metallicity estimates based on Lick indices for the old globular clusters of the M31 galaxy. The final aim is to add homogeneous spectroscopic metallicities to as many entries as possible of the Revised Bologna Catalog of M31 clusters\\thanks{RBC Version 4 available at: www.bo.astro.it/M31}, by reporting Lick indices measurements from any source (literature, new observations, etc.) into the same scale.} %methods {New empirical relations of [Fe/H] as a function of [MgFe] and Mg2 indices, as defined by Trager et al. (1998), are based on well studied Galactic Globular Clusters, complemented with theoretical model predictions for $-0.2\\le [Fe/H]\\le +0.5$. Lick indices for M31 clusters from various literature sources (225 clusters) and from new observations by our team (71 clusters) have been transformed into the Trager et al. (1998) system, yielding new metallicity estimates for  245 globular clusters of M31.  } %results {Our values are in good agreement with recent estimates based on detailed spectral fitting and with those obtained from Color Magnitude Diagrams of clusters imaged with the Hubble Space Telescope. The typical uncertainty on individual estimates is  $\\simeq \\pm 0.25$ dex, as resulted from the comparison with metallicities derived from Color Magnitude Diagrams of individual clusters.} %conclusions { The metallicity distribution of M31 globular cluster is briefly discussed and compared with that of the Milky Way. Simple parametric statistical tests suggest that the distribution is likely not unimodal.  The strong correlation between metallicity and kinematics found in previous studies is confirmed. The most metal-rich GCs tend to be packed at the center of the system and to cluster tightly around the galactic rotation curve defined by the HI disk, while the velocity dispersion about the curve increases with decreasing metallicity. However, also the clusters with $[Fe/H]<-1.0$ display a clear rotation pattern, at odds with their Milky Way counterparts.} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "introduction": "The concept of Simple Stellar Population (SSP) has proven to be a very fruitful tool for the study of virtually any kind of stellar system \\cite[][hereafter RFP88]{renbuz,renzfp}. A SSP is completely characterized by only four ``parameters'': (a)  mass, (b)  chemical composition, (c)  age, and (d) mass function, that determines the mass to light ratio (M/L) of the SSP once fixed the age and the chemical composition (see RFP88, for further possibly relevant variables that are not considered at zero-approximation). As a further simplification, the chemical composition is typically represented with two main parameters, i.e. the Helium abundance (Y) and the total abundance of elements heavier than He, usually parametrized by the iron abundance [Fe/H] (see, for instance \\cite{zwscale} (ZW84),\\cite{tc04mod,tc04alfa} and references therein). Even if the abundance of the so called $\\alpha$-elements has been subject of an increasing interest in the last two decades \\citep{mcw,marmod,tc04alfa,grat} [Fe/H] remains the main parameter to rank stars and/or stellar populations according to their abundance of heavy elements.  In the study of globular clusters (GC), which are the best approximation of a SSP in nature, the metallicity is a key parameter that is also needed to infer ages and age differences \\citep[see, for example RFP88 and][and references therein]{cg97}. The knowledge of the metallicity of a large sample of globular clusters in a given galaxy allows one to search for metallicity gradients, and the presence of distinct sub-populations of GCs (as in the Milky Way (MW) \\citet{Zinn}, and in many external galaxies \\citet{bs06}), and, in general, to obtain crucial information on the early phases of the formation and chemical enrichment of the parent galaxy.  While modern instrumentation has allowed the determination of the detailed abundance of several elements in single stars belonging to GCs of the MW \\cite[see][]{sneden,grat,euge}, the metallicities of extragalactic GCs must be obtained from their integrated colors and/or spectra, by comparison with Galactic templates and/or theoretical models \\citep{bs06}. Several broad-band integrated colors are fairly sensitive to metallicity and relatively easy to measure for clusters out to very large distances. However, they suffer from the well-known age-metallicity degeneracy (RFP88) {\\em and} they may be badly affected by the reddening due to extinction by interstellar dust. While the foreground extinction associated with the dust layers residing in our own Galaxy may be somehow constrained by observations and modeled \\citep{SFD}, the extinction intrinsic to external galaxies is largely unknown. On the other hand, spectral indices based on the strength of an absorption feature with respect to the surrounding continuum also suffer from the age-metallicity degeneracy \\citep[to different degrees, see][]{w94} but they are essentially unaffected by extinction \\citep{MacArthur}, a very desirable characteristic. The most widely used spectral indices were originally defined by \\cite{bur} and \\cite{faber} at the Lick Observatory. These authors defined a set of indices that measure the strength of the most pronounced absorption features that are seen in the integrated low-resolution spectra of stellar systems  at optical wavelengths. The use of Lick indices became widespread because they are easy to measure; as a consequence, they were also included as standard predictions in all theoretical models of SSP (see, for example, \\cite{buz_mg2,buz94,w94,bruz} hereafter BC03, \\citet{tc04mod,TMB}, hereafter TMB).  The M31 galaxy is an ideal target for studying GCs. It is nearby and it has a large cluster population ( $\\sim$3 times larger than the MW). The globular cluster system of M31 has been intensively studied in the past and several authors have used Lick indices to constrain the age and the metallicity of clusters in the Andromeda Galaxy.  Integrated-light spectroscopy of M31 GCs was pioneered by \\cite{syd69} who found that the GC system of this galaxy extends to higher metallicities with respect to the MW. In an important contribution, \\cite{bur} comprehensively discussed other interesting differences between GCs in M31 and in the MW. In particular they showed that M31 clusters have significantly stronger H$\\beta$ and CN absorption indices at a given metallicity. In a series of papers \\cite{bh90,bh91} and \\cite{huchra91} studied the metallicity distribution of M31 GCs using an extensive sample of integrated spectra and line indices. They found that the properties of the M31 GC system are broadly similar to the MW one, but they confirmed the presence of a high-metallicity tail having no counterpart in our Galaxy. They found that the mean metallicity  [Fe/H] was -1.2, and they identified a weak metallicity  gradient as a function of projected radius.  From the distribution of integrated colors \\cite{bar00} found  evidences that the M31 GC system may  have a bimodal metallicity distribution \\cite[like the Milky Way,][]{Zinn}, with peaks at [Fe/H]$\\sim -1.4$  and [Fe/H]$\\sim -0.6$.  Moreover, they found that the $(V-K)_0$ color  distribution was best modeled assuming three modes in the metallicity  distribution, instead of one or two.  Finally, they found a small radial metallicity gradient and no  correlation between cluster luminosity and metallicity  in M31 GCs. \\cite{perr_cat} produced a total sample of about 200  spectroscopic metallicities of M31 GCs, calibrating Lick indices measured in their own system versus the metallicity of M31 clusters in common with \\cite{huchra91} (hence they used a set of {\\em secondary} calibrators). They confirmed the bimodality in the metallicity  distribution and  reported that the metal--rich clusters have a higher rotation  amplitude with respect to metal-poor ones, while both groups are known to  rotate faster than their Galactic counterparts. Moreover, they found evidence for a radial metallicity gradient in the metal-poor population of M31 out to $\\sim 60\\arcmin$ from the galaxy center.  Metallicity (and age) estimates for various samples of M31 clusters obtained by fitting spectra with theoretical SSP models have been recently presented by  \\cite{beasII} and \\cite{puziam31}. \\cite{beasII} studied a sample of 23 M31 GCs with very high Signal to Noise (S/N) spectra,  seventeen of which were found to be old and to span a large range of metallicity, while the remaining six were classified as intermediate  age clusters. \\cite{puziam31} presented the metallicity of 70  globular clusters (including those studied by \\cite{beasII}) finding a bimodal distribution with peaks at [Z/H] $\\sim -1.66$ ($\\pm$0.05) and [Z/H] $\\sim -0.45$  ($\\pm$0.04) dex with dispersions 0.23 and 0.29 dex, respectively \\footnote{[Z/H] is defined as [Fe/H] + 0.94[$\\alpha$/Fe] taken from \\citet[][see also Trager et al. 2000]{TMB}.}.  More recently, \\cite{LeeII} merged the metallicities from \\cite{bar00} and \\cite{perr_cat}  with their own estimates from the  line indices measured from WIYN/Hydra spectra.  They found that a bimodal and trimodal distribution are statistically preferable to a unimodal metallicity distribution at a confidence level of 99.8\\%. \\cite{Fan} assembled metallicities from the literature and with estimates derived from integrated colors to obtain a global metallicity distribution of M31 GCs. They found a bimodal distributions with  peaks at [Fe/H]$\\sim$-1.7 and $\\sim$-0.7 dex with mean [Fe/H]=-1.21 dex, but  showed that three-group fits are also statistically acceptable. They  found a metallicity gradient as a function of projected radius for the metal-poor GCs, but no gradient for  the metal-rich GCs.  The brief summary of modern studies above underlines the great degree of heterogeneity of the available material. The various sets of estimates are obtained from observables that are different in nature (integrated spectral indices or colors) and are based on different calibrations (empirical, semi-empirical or theoretical; using primary or secondary calibrators). Even the actual definition of the same Lick indices varies from author to author; thus, in general, the presented calibrations are valid only for a given observational set-up and definition. In these circumstances it is clear that joining together different sets of metallicity estimates may be quite dangerous as it may lead to a poor degree of self-consistency in the final merged set.  We have assembled and we continuously maintain and update a database of parameters of confirmed and candidate clusters\\footnote{We also keep lists of targets previously suggested as candidate M31 globular clusters and later found to be objects of different nature, like distant galaxies, foreground stars, regions HII etc., see \\cite{2mass,io_rv,io_est}.} in M31, the Revised Bologna Catalog of M31 globular clusters \\cite[RBC,][]{2mass,io_rv,io_est}. As we want to add a reliable metallicity estimate to the confirmed clusters in the RBC, we need to devise the operational protocol to transform the actual measures provided in the literature (as, for example, already available and future sets of Lick indices for M31 clusters) into a unique homogeneous metallicity scale. In this paper, we describe the construction of this new homogeneous metallicity scale for M31 GCs based on Lick line indices. Having set and tested the new scale,  we present new metallicity estimates  for  245 M31 GCs.  The plan of the paper is as follows. In Section~\\ref{scale}, we describe the rationale and the procedure for the  construction of the new metallicity scale.  In Section~\\ref{data}, we report on the sample of M31 GC spectra that we have obtained and reduced, and from which we have estimated Lick indices. We also describe how we have collected Lick indices for M31 GCs from the literature and how we have reported all of them to the same homogeneous system. Finally, we derive the new values of the metallicities and, in Section~\\ref{comp}, compare  our scale with previous metallicity  estimates. In Section~\\ref{analysis}, we present and discuss the metallicity distribution of the M31 GC system and on the correlations between metallicity and kinematics. ",
        "conclusions": "}  Using our own data as well as datasets available in the literature, we have established a new homogeneous metallicity scale for M31 GCs. The scale is based on the Lick index Mg2 and on the combination of Mgb and Fe indices [MgFe], that have been calibrated against well studied Galactic globulars (for $[Fe/H]<-0.2$) and a variety of old-age SSP theoretical models for $-0.2 \\leq [Fe/H] \\leq 0.50$. Our scale has been shown to be self-consistent within $\\pm 0.25$ dex, and it should be applied only to classical, old globular clusters.  In the following we  briefly describe a few natural applications of the newly derived metallicity scale. In particular, (a) we derive and discuss the metallicity distribution of M31 GCs, (b) we explore the correlations between metallicity and kinematics,  for the sample of 240 bona-fide old GCs described in Sect.~3.4, above.  \\begin{figure*}[t] \\centering \\includegraphics[width=14cm, height=14cm]{AARfig17.ps} \\caption{Left Panels: Spatial distribution of three metallicity groups GCs in M31. The ellipses have a semimajor axis of 15, 30, 45, 60 arcmin. Right Panels:  Radial velocities vs. the projected distances along the major axis (X). The solid line shows a HI rotation curve from \\cite{hI}. } \\label{m31k3} \\end{figure*}  \\subsection{Metallicity distribution}  In Fig.~\\ref{m31md1} the Metallicity Distribution (MD) of our sample of M31 GCs is compared with its Milky Way counterpart.  The highest peak in the M31 MD occurs at [Fe/H]$\\sim -0.9$, coinciding with the overall average of the sample $<[Fe/H]> =-0.94$, significantly more metal rich than in the MW case, where the maximum is at [Fe/H]$\\sim -1.5$ and the overall mean is $<[Fe/H]>=-1.30$ \\cite[based on data from][that are in the ZW84 scale]{h96}. The M31 system appears also to have a much larger fraction of clusters having [Fe/H]$>-0.5$ (23\\% of the total sample) with respect to the Milky Way  (7\\%). It should be considered that the individual metallicity estimates for M31 clusters have much larger uncertainties with respect to their MW counterparts and this may produce some spurious widening of the MD for M31. However the shape of the distribution is essentially unchanged if we limit the analysis to the subset of clusters having errors in metallicity lower than $\\pm 0.3$ dex (132 clusters; dotted histogram in the upper panel of Fig.~\\ref{m31md1}). Fig.~\\ref{m31md2} shows that the difference between the MDs of the two galaxies cannot be ascribed to spurious effects due to our calibration, as it can be re-conduced to genuine differences in the observable [MgFe].  The MD of M31 GCs do not present any obvious structure like the bimodality encountered in the GC system of the Milky Way. Nevertheless the distribution for M31 clusters does not seem to be well represented by a single Gaussian distribution. Note that large errors on individual metallicities should contribute to wipe out real structures, not to produce spurious ones. The hypothesis of a multimodal underlying distribution has been compared with a unimodal representation using the parametric KMM test \\citep{KMM}, that compares the fits to the MD made with one or more Gaussian distributions. A two component model with modes at [Fe/H]=-1.54 and [Fe/H]=-0.64 is preferred to the unimodal case at the 99.1\\% confidence level (homoscedatic case) and at the 98.7\\% c.l. in the  heteroscedastic case with peaks at [Fe/H]=-1.79 and [Fe/H]=-0.76.  A three component model with modes at $[Fe/H]=-0.25$,  $-0.89$ and $-1.72$ is also preferable to the unimodal one (99.8\\% c.l., in the homoscedastic case, and 99.6\\% c.l., in the heteroscedastic case with peaks at [Fe/H]=-1.77, [Fe/H]=-0.80 and [Fe/H]=-0.01), and nearly equivalent to the bimodal representation, from a statistical point of view. The preference of bi- and three- modal models over the unimodal case remains even if we  consider the subset of clusters with the lowest metallicity errors described above. While clearly not conclusive, the above analysis suggests that there may be real structures in the MD of M31 globular clusters, in good agreement with the conclusions reached by \\citet{bar00}, P02, Pz05, \\cite{Fan} and \\cite{LeeII}.  \\subsection{Metallicity and kinematics}  Fig.~\\ref{m31k3} shows the positional and kinematical properties of M31 GCs divided into three groups according to their metallicity, i.e. a Metal Poor (MP) group ($[Fe/H]\\le -1.0$), a Metal Intermediate (MI) group ($-1.0<[Fe/H]<-0.5$), and a Metal Rich (MR) group ($[Fe/H]\\ge -0.5$). The left panels of Fig.~\\ref{m31k3} show the spatial distribution of the considered clusters in the canonical X,Y projected coordinate system \\cite[see][and references therein]{2mass}, with X along the major axis of the galaxy. In the right panels the radial velocity of the clusters (in the reference frame of M31) is plotted versus the X coordinate and compared with the rotation curve of the HI disk from \\cite{hI}.  It results quite clear from the inspection of Fig.~\\ref{m31k3} that the MR and MI subsamples display a significant rotation pattern, much similar to the rotation curve of neutral Hydrogen disk of M31. The MR clusters are more densely packed near the center of the galaxy and appear to follow more closely the HI curve, whereas the MI clusters display a larger dispersion. MR clusters are likely associated with the prominent Bulge of M31.  The MP clusters show a much larger velocity dispersion at any distance from the center of the galaxy; in spite of that, they follow a significant rotation pattern in the same sense as the other clusters. Dividing the MP sample at X=0, we find (M31-centric) average velocities of $\\langle V_{M31}\\rangle=+59$ km/s and $\\langle V_{M31}\\rangle=-48$ km/s for the clusters with $X>0$ and $X<0$, respectively; the difference in the median velocities is even larger, as $V_{M31}^{med}=+86$ km/s and $V_{M31}^{med}=-59$ km/s, for the two subsets. Finally if the $V_{M31}$ distributions of the $X>0$ and $X<0$ MP clusters are compared with a Kolmogorov-Smirnov test, it turns out that the probability that the two samples are drawn from the same distribution of $V_{M31}$ is just 0.2\\%. In Sect.~3.4 we have shown that our sample should be reasonably clean from spurious sources \\cite[as for instance young massive clusters, that may be misclassified as metal-poor GCs and would follow the rotation pattern of the thin disk they belong to, see][]{blcc}, hence we conclude that the rotation pattern of MP clusters is probably real. However an ultimate conclusion on this (relevant) issue could be achieved only when the actual nature of a significant subsample of MP clusters will be confirmed beyond any doubt from the CMD of their individual stars.  The above results are in good agreement with what previously found by P02 and \\cite{LeeII}, among  others.  A more detailed discussion of these correlations between kinematics and metallicity is beyond the scope of the present paper. We address the interested reader to the thorough discussion by \\cite{LeeII}. Here we just want to draw the attention of the reader on five of the six MR clusters lying at $R>30\\arcmin$ (labeled in Fig.~\\ref{m31k3}). B001, B398 and B403 show no correlation with the overall rotation pattern.  On the other hand B292D, and B457  lie straight  on the flat branch of the HI rotation curve in spite of the fact that they are more than $\\sim 7$ kpc away from all the other MR clusters (except B398 and B403, of course).  These five objects clearly deserve new observations with high S/N spectra to verify both their metallicity and their radial velocity. If confirmed, their odd positions and kinematics would require an interpretation. Moreover, B403 and B407 (labeled in the MI panel of Fig.~\\ref{m31k3}, and having  $[Fe/H]=-0.65\\pm 0.15$) have very similar position and velocity (differing by  $\\sim20$~km/s). The case of these two relatively metal-rich clusters in the outer halo of M31 is discussed in more detail in \\cite{Perinacmd}.\\footnote{{\\em Note Added in Proofs:} After the acceptance of the present paper, a work appeared presenting metallicities of a few M31 GCs from high resolution integrated spectra \\citep{colucci}. It is interesting to note that for four of the five clusters in common, our [Fe/H] estimates agree with those of \\citet{colucci} to within $<0.1$ dex. Our estimate for the fifth cluster (B358, which is in the deep metal-poor regime in which our observable are less sensitive to metallicity, see Sect.~2, above) is 0.36 dex higher than what found by \\citet{colucci}, i.e. $[Fe/H]=-1.85\\pm 0.19$ vs.\\ $[Fe/H]=-2.21\\pm 0.03$; however the difference is less than $2\\sigma$. It appears that our own scale and that based on high-resolution spectroscopy introduced by these authors are in {\\em excellent} agreement. We note also that our metallicity estimate for the remote M31 cluster MGC1, $[Fe/H]=-2.16\\pm 0.28$, is in excellent agreement with what recently obtained by \\citet{macGC1} from a high-quality CMD, $[M/H]\\simeq -2.3$. }"
    },
    "0910/0910.1661_arXiv.txt": {
        "abstract": "The so-called \\textsf{solar cycle} is generally characterized by the quasi-periodic oscillatory evolution of the photospheric spots number. This quasi-periodic pattern has always been an intriguing question. Several physical models were proposed to explain this evolution and many mathematical data analysis were employed to determine the principal frequencies noticeable in the measured data. Both approaches try to predict the future evolution of the solar activity and to understand the physical phenomena producing these cycles. Here we present the analysis of the sunspots number evolution using the time-delay approach. Our results show than the solar cycle can also be characterized by this behavior implying the influence of the past evolution over the present one, suggesting an histeresis mechanism, linked probably with magnetic activity. ",
        "introduction": "Solar activity can be seen through the evolution of sunspots number in quasi-periodic oscillatory series with periods going from 8 to 15 years and with a mean period of 11~years. Due to the change of magnetic field polarity in solar hemispheres alternatively each cycle, the period is rather 22 years. The quasi-periodic evolution of % many activity phenomena % is still a unsolved key problem in solar physics (along with, e.g., heating of the solar chromosphere and corona, and solar flares). An important issue of this understanding is due to the influence of solar activity over the terrestrial climate (\\cite[Archibald 2006]{arch06}). %  \\medskip Many physical models were proposed to understand the basic mechanisms (e.g., \\cite[Benevolenskaya 1998]{bene98}) producing solar/stellar activity and its quasi-cyclic evolution. % From sunspots time series, mathematical approaches have highlighted several other hidden periods other that the well-apparent 11-years period. These works contributed to provide some observed parameters to constraint theoretical models, and to predict the future evolution of the sunspots number (e.g., \\cite[Clilverd \\etal\\ 2003]{clil03}; \\cite[Sello 2003]{sello03}). %  \\medskip In this work we follow another mathematical approach, investigating the influence of the past evolution of the sunspots number over the present one. A natural and simple way to carry out this analysis is to assume that the susnpots number have a temporal delay behaviour. Thus, the relationship between present and past values becomes non-linear and could prove that the past amplitude influence the present one. As for the abovementioned methods, this one also allows some predictions about future solar cycles by using intermediate parameters to characterize the general past evolution.  \\vspace*{-0.3 cm} ",
        "conclusions": "The $T$ values which minimize the correlation between the sunspots number ratio and the past spots number, are 7 or 8 years (Fig.~\\ref{Fig:RatioFactor}). To test the accuracy of our method, we applied it on the solar cycle~23 (Table~\\ref{Tab:Cycle23} and Fig.~\\ref{Fig:EvolSpots}, left panel). Our predictions are consistent with the observations, except for the unusually long period of low activity at the end of this cycle. We also characterized the next solar cycle~$24$ (Table~\\ref{Tab:Prediction} and Fig.~\\ref{Fig:EvolSpots}) using our modelling of the solar cycle~$23$ as well as the spots number observed until $2006.5$ and $2007.5$. The next maximum sunspots number would take place between $2011$ and $2012$ and should be close to 60 (Fig.~\\ref{Fig:EvolSpots}, right panel) as for the solar cycle~$14$. Using the solar minimum occurred in $2008.5$ ($R_{min} = 2.9$ from SDIC), the predicted maximum is consistent with the value from the linear least-square fit on the both axis (red line on Fig.~\\ref{Fig:RminRmax}) and is a few smaller than the range given by \\cite[Brajsa \\etal\\ (2009)]{Brajsa09}. Moreover the next solar minimum should not occur before 2019 or even 2020.   \\begin{table}[h] \\begin{center} \\caption{ {\\scriptsize Observed and predicted values for the cycle 23. Results showed in Fig.~\\ref{Fig:EvolSpots} (left panel).}} \\label{Tab:Cycle23} {\\scriptsize \\begin{tabular}{cccccccc} \\noalign{\\medskip} \\hline \\multicolumn{2}{l}{Observed parameters} & \\multicolumn{2}{c}{Epoch of solar} & \\multicolumn{2}{c}{Maximum sunspot} & \\multicolumn{2}{c}{Epoch of the end} \\\\ \\multicolumn{2}{l}{of solar minimum}       & \\multicolumn{2}{c}{maximum}        & \\multicolumn{2}{c}{number}                 & \\multicolumn{2}{c}{of cycle}              \\\\ \\noalign{\\medskip} Epoch                & Number                    & Observed        & Predicted            & Observed        & Predicted                & Observed        & Predicted               \\\\ \\hline $1996.4$            & $8.0$                        & $2000.3$        & $2001.5$            & $120.8$           & $106.4$                   & $2009.?$        & $2006.5$                \\\\ \\hline \\noalign{\\medskip} \\end{tabular} } \\end{center} \\end{table}  \\begin{figure}[h] \\begin{center} \\includegraphics[width=1.6in]{Solar_Cycle_ss_label_1900-1997_2020_Dec_7.eps} \\includegraphics[width=1.55in]{Solar_Cycle_ss_label_1900-2007_2020_Dec_8.eps} \\includegraphics[width=1.55in]{Solar_Cycle_ss_label_1900-2008_2020_Dec_8.eps} \\caption{\\scriptsize The observed (\\textit{plus symbols and dotted line}) and selected (\\textit{asterisks and blue line}) yearly values of the relative sunspots number. Our predictions using the temporal delay method are also marked (\\textit{diamonds and red line}). \\textit{Left panel:} Prediction for the solar cycle 23 and 24. \\textit{Middle and Right panels:} Prediction of the solar cycle~24 using the sunspots number observed until 2006.5 and 2007.5 as input.} \\label{Fig:EvolSpots} \\end{center} \\end{figure}  \\begin{table}[h] \\begin{center} \\caption{\\scriptsize Predicted values for the solar cycle 24 using various data set in input.} \\label{Tab:Prediction} {\\scriptsize \\begin{tabular}{lcccccc} \\noalign{\\medskip} \\hline Panel of                        & \\multicolumn{2}{c}{Input parameters of} & \\multicolumn{2}{c}{Predicted parameters} & \\multicolumn{2}{c}{Predicted parameters of} \\\\ Fig.~\\ref{Fig:EvolSpots} & \\multicolumn{2}{c}{solar minimum}        & \\multicolumn{2}{c}{of solar maximum}      & \\multicolumn{2}{c}{next solar minimum}  \\\\ \\noalign{\\medskip} & Epoch                & Number                      & Epoch                & Number                    & Epoch                & Number  \\\\ \\hline Left        & $2006.5$            & $19.6$                        & $2011.5$            & $~~84.1$                 & $2016.5$             & $23.0$  \\\\ Middle    & $2006.5$            & $15.2$                        & $2012.5$            & $137.3$                   & $2018.5$             & $28.8$  \\\\ Right      & $2007.7$            & $~~7.5$                      & $2012.5$            & $~~57.1$                 & $2020.5$             & $~~4.2$  \\\\ \\hline \\noalign{\\medskip} \\end{tabular} } \\end{center} \\end{table}  \\begin{figure}[h] \\vspace*{-0.6 cm} \\begin{minipage}{0.65\\textwidth} \\caption{\\scriptsize Maximum ($R_{max}$) versus minimum ($R_{min}$) sunspots number of a given solar cycle (\\textit{dots}). The cycle number is also marked. Blue, green and red diamonds show the locus of our predicted maxima using our modelling of the solar cycle~$23$ as well as the spots number observed until $2006.5$ and $2007.5$, respectively (Table~\\ref{Tab:Prediction}). The predicted range of \\cite[Brajsa \\etal\\ (2009)]{Brajsa09} is plotted as well (\\textit{vertical line}). The blue and red lines show linear fits using data of both axis and the linear least-square fit like one given by \\cite[Brajsa \\etal\\ (2009, Fig.~6)]{Brajsa09}, respectively.} \\end{minipage} \\hfill \\begin{minipage}{0.35\\textwidth}  \\vspace*{0.35 cm} \\centerline{ \\includegraphics[width= 1.6in]{Fig_Rmin_Rmax_mod.ps}} \\label{Fig:RminRmax} \\end{minipage} \\end{figure}   These preliminary results are encouraging because we find a similar delay as that observed between the geomagnetic activity and solar cycle peaks \\cite[(Hathaway \\& Wilson 2006)]{hat06} linked probably with magnetic activity by some kind of histeresis mechanism. At present, we can successfully reproduce some previous solar cycles (e.g., epoch and solar maximum). We plan to improve our predictive method including the influence of delays for which there is a good correlation. This more detailed description could allow to solve the overestimation of the next solar minimum and to reproduce the asymmetries observed during the previous cycles. Our final aim is to obtain a reliable prediction of the whole solar activity by identifying the fundamental lower and high activity precursors of its present and future evolution. All these improvements are still needed to be able to predict future low solar activity periods (e.g., Maunder/Dalton Minimum)."
    },
    "0910/0910.4272_arXiv.txt": {
        "abstract": "{The rotation curve, the total mass and the gravitational potential of the Galaxy are sensitive measurements of the dark matter halo profile.} {In this publication cuspy and cored DM halo profiles are analysed with respect to recent astronomical constraints in order to constrain the shape of the Galactic DM halo and the local DM density. } {All Galactic density components (luminous matter and DM) are parametrized. Then the total density distribution is constrained by astronomical observations: 1) the total mass of the Galaxy, 2) the total matter density at the position of the Sun, 3) the surface density of the visible matter, 4) the surface density of the total matter in the vicinity of the Sun, 5) the rotation speed  of the Sun and 6) the shape of the velocity distribution within and above the Galactic disc. The mass model of the Galaxy  is mainly constrained by the local matter density (Oort limit), the rotation speed of the Sun and the total mass of the Galaxy from tracer stars in the halo. } {It is shown from a statistical $\\chi^2$ fit to all data that the local DM density is strongly positively (negatively) correlated with the scale length of the DM halo (baryonic disc). Since these scale lengths are poorly constrained the local DM density can vary from 0.2 to 0.4 GeV cm$^{-3}$ (0.005 - 0.01 M$_\\odot$ pc$^{-3}$) for a spherical DM halo profile and allowing total Galaxy masses up to 2 $\\cdot$ 10$^{\\mathrm{12}}$ M$_\\odot$. For oblate DM halos and dark matter discs, as predicted in recent N-body simulations, the local DM density can be increased significantly. } {} ",
        "introduction": "\\label{sec:Introduction}  The best evidence for Dark Matter (DM) in galaxies is usually provided by rotation curves, which do not fall off fast enough at large distances from the centre. This can  be understood, either by assuming Newton's law of gravitation is not valid at large distances, the so-called MOND (Modified Newtonian Dynamics) theory \\citep{Bekenstein:1984tv,Bienayme:2009wb}, or the visible mass distribution is augmented by invisible mass, i.e. DM. For a flat rotation curve the DM density has to fall off like 1/r$^2$ at large distances.  Given the overwhelming evidence for DM on all scales from the flatness of the universe combined with gravitational lensing and structure formation we assume DM exists and try to constrain the DM density from dynamical constraints, for which better data became available in recent years: \\begin{enumerate} \\item the total mass of the Galaxy has to be about 10$^{12}$ solar masses \\citep{Wilkinson:1999hf, Battaglia:2005rj,Xue:2008se}; \\item the total mass inside the solar orbit is constrained by the well-known rotation speed of the solar system; \\item the total matter density at the position of the Sun from the gravitational potential determined from the movements of local stars as measured with the Hipparcos satellite \\citep{Holmberg:2004fj}; \\item the surface density of the visible matter at the position of the Sun \\citep[and references therein]{Naab:2005km}; \\item the surface density of the total matter at the position of the Sun \\citep{Kuijken:1991mw, Holmberg:2004fj,Bienayme:2005py}; \\item the shape of the rotation curve within Galactic disc \\citep[and references therein]{Sofue:2008wt}; \\item the  velocity distribution above the Galactic disc ($z>4$ kpc) \\citep{Xue:2008se}. \\end{enumerate} Towards the Galactic centre (GC) N-body simulations of structure formation predict a DM density, which diverges at the centre, thus developing a so-called cusp, with the slope varying more like 1/r than 1/r$^2$ \\citep{Navarro:1996gj,Moore:1999nt,Ricotti:2002qu}. On the contrary,  rotation curves of nearby dwarf galaxies do not show such a cusp, but rather a constant density near the center, a so-called core \\citep{Oh:2008ww,Gentile:2007sb,Salucci:2007tm}. N-body simulations including the complicated baryonic disc formation indicated that the cusps may have been destroyed by the fluctuation in the gravitational potential of the baryons, which is much stronger than the potential caused by the DM in the centre \\citep{Mashchenko:2006dm}. In addition, the DM may form a disc like structure by the infall of DM dwarf galaxies, as shown by recent high resolution N-body simulations \\citep{Purcell:2009yp}. Evidence for substructure of DM inside the disc is supported by the structure in the gas flaring \\citep{Kalberla:2007sr}, by the peculiar change of slope for the rotation curve inside the disc, which is not observed for stars outside the disc (Sect. \\ref{subsec:RC}) and by the structure in the diffuse gamma radiation \\citep{deBoer:2005tm}.  Such substructure will not be investigated in this paper, but smooth DM halos with different (cored and cuspy) profiles will be compared with all available data. These will result in lower limits on the local DM density, since  dark matter discs or other local substructure will only enhance the local density.  A reliable determination of the local DM density is of great interest for direct DM search experiments, where elastic collisions between WIMPs and the target material of the detector are searched for. This signal is proportional to the local density. A review on direct searches can be found in the paper by \\citet{Spooner:2007zh}.  The structure of the paper is as follows:  in section \\ref{sec:para} the parametrization of the luminous matter and five different DM halo profiles are given. In Sect. \\ref{sec:data} the experimental data used to determine the mass model are discussed. Then the numerical determination of the mass model parameters using a $\\chi^2$ fit of all dynamical constraints is discussed in Sect. \\ref{sec:fit}. At the end a summary of the results is given. ",
        "conclusions": "\\label{sec:Conclusion} In this analysis five different halo profiles are compared with recent dynamical constraints as summarized in Table \\ref{param}. The change of slope in the RC around 10 kpc (Fig. \\ref{fig:rc_simple}) was ignored, so the monotonical decreasing RC for the smooth halo profiles do not describe the data well. The change of slope may be related to a ringlike DM substructure , as indicated by the structure in the gas flaring \\citep{Kalberla:2007sr} and by the structure in the diffuse gamma radiation \\citep{deBoer:2005tm}. Such a ringlike structure of DM gives a perfect description of the rotation curve, especially the fast decrease between 6 and 10 kpc. If the DM substructure is included, the local DM density increases above the values found in this analysis, so the  values quoted here should be considered lower limits.  The astronomical constraints are consistent with a density model of the Galaxy consisting of a central bulge, a disc and an extended DM halo with a cuspy density profile and a local DM density between 0.2 GeV cm$^{-3}$ (0.005 M$_\\odot$ pc$^{-3}$) and 0.4 GeV cm$^{-3}$ (0.01 M$_\\odot$ pc$^{-3}$), as shown in Fig. \\ref{fig:diffuse_simple}. Strong positive and negative correlations between the parameters were found in the fit and they are causing the obvious correlations between $\\rho_{\\odot,\\mathrm{DM}}$ and M$_{\\mathrm{tot}}$ in Fig. \\ref{fig:diffuse_simple}. For non-spherical haloes these values can be enhanced by 20\\%. If dark discs are considered, densities up to 0.7 GeV cm$^{-3}$ (0.018 M$_\\odot$ pc$^{-3}$) can be easily imagined, so the previous quoted range of 0.2 - 0.7 GeV cm$^{-3}$ (0.005 - 0.018 M$_\\odot$ pc$^{-3}$) seems still valid. This range is considerably larger than the values quoted by analyses which used a Markov Chain method to minimize the likelihood; they find $\\rho_{\\odot,\\mathrm{DM}} = 0.39 \\pm 0.03$ GeV cm$^{-3}$ \\citep{Catena:2009mf} and $\\rho_{\\odot,\\mathrm{DM}} = 0.32 \\pm 0.07$ GeV cm$^{-3}$ \\citep{Strigari:2009zb} respectively. But given the good $\\chi^2$ values for our fits obtained for a large range of DM densities we see no way that the errors can be as small as quoted by these authors."
    },
    "0910/0910.1609.txt": {
        "abstract": " ",
        "introduction": "\\label{sec:intro}  The Atacama Large Millimeter/submillimeter Array (ALMA) will be a single research instrument composed of at least fifty-four 12-m and twelve 7-m high-precision antennas, located at a very high altitude of 5000 m on the Chajnantor plain of the Chilean Andes.  The weather conditions at the ALMA site will allow transformational research into the physics of the cold Universe across a wide range of wavelengths, from radio to submillimeter.  Thus, ALMA will be capable of probing the first stars and galaxies and directly imaging the disks in which planets are formed. ALMA will be a complete astronomical imaging and spectroscopic instrument for the millimeter/submillimeter regime, providing scientists with capabilities and wavelength coverage that complement those of other research facilities of its era, such as the Expanded Very Large Array (EVLA), James Webb Space Telescope (JWST), Thirty Meter Telescope (TMT), European Extremely Large Telescope (E-ELT),   and Square Kilometer Array (SKA). ALMA will be the preeminent platform for astronomical research at millimeter (mm) and submillimeter (submm) wavelengths for decades to come.  ALMA will revolutionize many areas of astronomy and an amazing breadth of science has already been proposed (see for example the ALMA Design Reference Science Plan). The technical requirements of the ALMA Project are, however, driven by three specific Level One Science Goals:\\\\ \\noindent{\\bf (1)} The ability to detect spectral line emission from CO or C+ in a normal galaxy like the Milky Way at a redshift of $z = 3$, in less than 24 hours of observation.\\\\ \\noindent{\\bf (2)} The ability to image the gas kinematics in a solar-mass protostellar/ protoplanetary disk at a distance of 150 pc (roughly, the distance of the star-forming clouds in Ophiuchus or Corona Australis), enabling one to study the physical, chemical, and magnetic field structure of the disk and to detect the tidal gaps created by planets undergoing formation.\\\\ \\noindent{\\bf (3)} The ability to provide precise images at an angular resolution of 0.1 arcseconds. Here the term ``precise image\" means an accurate representation of the sky brightness at all points where the brightness is greater than 0.1\\% of the peak image brightness. This requirement applies to all sources visible to ALMA that transit at an elevation greater than 20\\degree.  ALMA was originally envisioned to provide access to all frequencies between 31 and 950 GHz  accessible from the ground.  During the re-baselining exercise undertaken in 2001, the entire project was scrutinized to find necessary cost savings and the two lowest receiver frequencies, covering 31--84 GHz, were among those items delayed beyond the start of science operations. Nevertheless, Band 1, covering 31--45 GHz, was re-affirmed as being a high priority future item for ALMA.  In May 2001, John Richer and Geoff Blake prepared a document {\\it Science with Band 1 (31--45 GHz) on ALMA} as part of the re-baselining exercise. Key arguments for maintaining the receiver were put forward and included: (1) exciting science opportunities, bringing in a wider community of users; (2) a significantly faster imaging and survey instrument than the upgraded EVLA, especially due to the larger primary beam; (3) access to the southern sky at these wavelengths; (4) excellent science possible even in ``poor\" weather; (5) a relatively cheap and reliable receiver to build and maintain; and (6) a very useful engineering/debugging tool for the entire array.  The Richer/Blake document was followed by an ASAC Committee Report in October 2001 in which the addition of Japan into the ALMA project re-opened the question of observing frequency priorities for those receiver bands which had been put on hold during re-baselining. The unanimous recommendation of the ASAC was to put Band 10 as top priority, followed by a very high priority Band 1. The key science cases for Band 1, at that time, were seen to be (1) high resolution Sunyaev-Zel'dovich (SZ) imaging of cluster gas at all redshifts; and (2) mapping the cold ISM in Galaxies at intermediate and high redshift.  The scientific landscape has changed significantly since 2001 and thus it is time to update the main science drivers for Band 1. The ALMA Development Plan also now lies on the horizon, and those parties interested in longer wavelength observing with ALMA recognize that now is the time to put forward their best cases. Thus, in the Fall of 2008, two dozen astronomers from around the globe met in Victoria, B.C. to discuss this issue.  This paper represents both a record of that meeting and an attempt to move toward concrete proposals for outstanding science with Band 1. In Section \\ref{sec:one} we present two science cases that reaffirm and enhance the already established ALMA Project Level One Science Goals. In Section \\ref{sec:range} we provide a selection of science cases that reinforce the breadth and versatility of the Band 1 Receiver. Section \\ref{sec:consid} discusses both weather considerations at the ALMA site and the complementarity of ALMA versus the EVLA at Band 1 frequencies. Finally, Section \\ref{sec:conc} summarizes the paper. ",
        "conclusions": "\\label{sec:conc}  The 31--45 GHz Band 1 Receiver has been considered an essential part of ALMA from the earliest planning days. Even through the re-baselining exercise in 2001, the importance of Band 1 was emphasized. With the ALMA Development Plan taking shape, we have undertaken an updated review of the scientific opportunity at these longer wavelengths. This document presents a set of compelling science cases for Band 1. The science ranges from nearby stars and galaxies to the re-ionization edge of the Universe.  Two of the science cases provide additional leverage on the present ALMA Level One Science Goals and are seen as particularly powerful motivations for building the Band 1 Receiver: (1) detailing the evolution of grains in protoplanetary disks, as a complement to the gas kinematics, requires continuum observations out to $\\sim 30\\,$GHz ($\\sim 1\\,$cm); (2) detecting CO 3 -- 2 spectral line emission from Galaxies like the Milky Way during the era of re-ionization, $6.5 < z < 10$ also requires Band 1 Receiver coverage.   %"
    },
    "0910/0910.0058_arXiv.txt": {
        "abstract": "The [Ne~{\\sc ii}] fine-structure emission line at 12.8$\\mu$m has been detected in several young stellar objects (YSO) spectra. This line is thought to be produced by X-ray irradiation of the warm protoplanetary disk atmospheres, however the observational correlation between [Ne~{\\sc ii}] luminosities and measured X-ray luminosities shows a large scatter. Such spread limits the utility of this line as a probe of the gaseous phase of disks, as several authors have suggested pollution by outflows as a probable cause of the observed scatter. In this work we explore the possibility that the large variations in the observed [Ne~{\\sc ii}] luminosity may be caused instead by different star-disk parameters. In particular we study the effects that the hardness of the irradiating source and the structure (flaring) of the disk have on the luminosity and spectral profile of the [Ne~{\\sc ii}]~12.8 $\\mu$m line. We find that varying these parameter can indeed cause up to an order of magnitude variation in the emission luminosities which may explain the scatter observed, although our models predict somewhat smaller luminosities than those recently reported by other authors who observed the line with the {\\it Spitzer} Space Telescope. Our models also show that the hardness of the spectrum has only a limited (undetectable) effect on the line profiles, while changes in the flaring power of the disk significantly affect the size of the [Ne~{\\sc ii}] emission region and, as a consequence, its line profile. In particular we suggest that broad line profiles centred on the stellar radial velocity may be indicative of flat disks seen at large inclination angles. ",
        "introduction": "\\indent Young solar mass stars are surrounded by $\\sim10^{-2}$M$_{\\odot}$ of dust and gas distributed in a circumstellar disk which is the birthplace of planetary systems. In recent years efforts have been made to study the inner regions ($\\la 10$~AU) of these disks in order to constrain gaseous planetary system theories, and to understand the physical processes that characterise the warm and chemically active regions of protoplanetary disk atmospheres. Accurate theoretical modelling of the spectral energy distribution (SED) of young stars (see reviews by Dullemond et al. 2007 and Natta et al. 2007), coupled with observations in the near and mid-infrared carried out with the $Spitzer$ Space Telescope have greatly improved our understanding of the dust component in the inner disks. On the other hand, direct studies of the gaseous component of these disks are more challenging, even though the gas exceeds the dust by over two orders of magnitude in mass. So far molecular line emission, like the CO fundamental and overtones \\citep{Najita2007}, the rovibrational transition of water \\citep{Carr2008,Salyk2008},  and H$_{2}$ \\citep{Herczeg2002}, have been used as diagnostics of regions with temperatures around 1000-2000~K.  The [Ne~{\\sc ii}]~12.8$\\mu$m line (from now on [Ne~{\\sc ii}]) has been suggested as a new probe of gas in the potentially planet forming regions of the inner disks (e.g. Glassgold et al. 2007, Hollenbach \\& Gorti, 2009). The ionised region of the disk where this line is formed is heated mainly by X-rays from the central star (Glassgold et al. 2007, Meijerink et al. 2008, Ercolano et al. 2008a, Glassgold, Ercolano \\& Drake, 2009; but see also Gorti \\& Hollenbach, 2008). Indeed, observations from space, with the {\\it Spitzer} Infrared Spectrometer (IRS) \\citep{Pascucci2007,Lahuis2007,Espaillat2007}, and from the ground, with MICHELLE and TEXES at Gemini North (Herczeg et al., 2007; Najita et al., 2009) and VISIR at VLT (Van Boekel et al. 2009), have successfully detected the [Ne~{\\sc ii}] line in tens of young systems. The mean observed line luminosity is comparable to the predicted value from disk models, but the scatter of the observed values around this mean spans over three orders of magnitude (e.g. G\\\"udel et al. 2009). If the [Ne~{\\sc ii}] line is to be used as a useful diagnostic of the gaseous phase of the inner disks its origins must be well understood. In particular if the line is formed in the disk by irradiation of the gas by X-rays from the central star then one naturally expects to find a correlation between line luminosity and X-ray luminosity for the observed sources. Such correlation is indeed observed (Pascucci et al 2007, G\\\"udel et al. 2009), however, for objects with the same X-ray luminosity the [Ne~{\\sc ii}] line luminosities span over at least one order of magnitude even after removing those sources with known outflows, for which it can be argued that a large contribution to the line comes from the outflow itself (e.g. T Tau, van Boekel et al. 2009). Understanding the origin of this scatter is extremely important if useful information is to be extracted from future observations of this line. For example, one should expect that disk structural properties, like presence of holes or different degree of flaring, not considered in earlier models have an effect on the  [Ne~{\\sc ii}]  emitting region. Moreover, one relevant question is whether the [Ne~{\\sc ii}] line originates in the bound layers of the disk irradiated mainly by X-rays or whether it includes a contribution from a photoevaporative outflow (e.g. Alexander, 2008). In this paper we will attempt to answer the latter question by exploring the possibility that the scatter observed in [Ne~{\\sc ii}] line luminosities for a given X-ray luminosity may be due to details of the irradiating field or to the disk structure (i.e. degree of flaring). We will also produce theoretical line profiles and will compare them to available high resolution observations (e.g. Herczeg et al. 2007, Najita et al. 2009).  We use the models of Ercolano et al. (2009) to calculate [Ne~{\\sc ii}] line luminosities and line profiles for irradiating spectra of varying hardness. We also show the results from additional models that were run to investigate the effects of disk flaring on the line luminosities and profiles. The paper is organised as follows. The modeling strategy is outlined in Section 2; Section 3 contains a review of our results, while our final conclusions and discussion are given in Section~4. ",
        "conclusions": "As the sample of observed L([Ne~{\\sc ii}]) luminosities from protoplanetary disks continues to grow, together with the sample of pre-main sequence stars with measured X-ray luminosities, a trend between the two quantities becomes more and more evident (G\\\"udel et al. 2009). While this lends evidence to the [Ne~{\\sc ii}] line being a product of X-ray irradiation of disks as predicted by Glassgold et al. (2007), the large scatter associated with this relation makes the situation less clear-cut. For a handful of objects with known powerful outflows it is clear that this line cannot be used as a diagnostics of the gaseous {\\it disk} phase, as the emissivities may be dominated by gas in the outflows. However, even after removing objects with known outflows from our sample we are left with approximately {\\it one order of magnitude variations in L([Ne~{\\sc ii}]) at any given L$_X$}.  In this paper we have investigated the origin of the large scatter observed in the [Ne~{\\sc ii}] -- L$_X$ relation. We find that variations in the irradiating spectral shape and disk structure (flaring) are sufficient to explain the typically observed scatter of approximately one order of magnitude. Figure~\\ref{Fig:LNeIILX} shows the results from our models (coloured symbols) overplotted on a shaded region representing roughly the observations presented by G\\\"udel et al. (2009), excluding the strong outflow objects (blue points in their plot).  The definition of such region is affected by the limited sample statistic that does not allow a sharper outline, but it shows where the bulk of data falls. Moreover, we note here that we use luminosities {\\it after} the screen while G\\\"udel et al. use intrinsic luminosities derived from the observations. We have performed absorption calculations for a subset of the observations from which we found the following: absorbing the intrinsic X-ray spectral models by $N_{\\rm H} = 3\\times 10^{20}~{\\rm cm}^{-2} - 4.6\\times 10^{22}~{\\rm cm}^{-2}$ (a range occupied by the observations) reduces the 0.1-10~keV luminosity by factors of approximately 1.4--5.5 (also depending on the intrinsic thermal model from the individual fits). On average, thus, the {\\it unabsorbed} $L_{\\rm X}$ from the observations should be shifted  by a factor of 3 (0.5~dex) to the left to obtain average absorbed values, with individual shifts ranging from 0.15~dex to 0.75~dex. These corrections do not change the plot  qualitatively. We, however, refrain from using the individual {\\it absorbed} (post-screen) luminosities derived from the observations here because it is not clear a priori that the absorbed luminosity seen from Earth is close to the absorbed luminosity seen by the disk. The latter luminosity is principally unknown. We also mention that G\\\"udel et al. use the 0.3-10~keV range for their luminosities, while we use 0.1-10~keV. For the post-screen luminosities for stars with $N_{\\rm H} > 10^{20}$~cm$^{-2}$, however, the difference is minor.  Our model predictions seem to agree well with the observations in terms of reproducing the scatter, and at intermediate and high X-ray luminosities the absolute values of L([Ne~{\\sc ii}]) also roughly agree with the observations. Models with low X-ray luminosities, however, fall short of the observed values, and while detection limits can partially affect the observations, the failure of our models to produce L([Ne~{\\sc ii}]) $>$ 10$^{28}erg/sec$ for X-ray luminosities of 2$\\times$10$^{29}erg\\,s^{-1}$ must have another origin. Nevertheless there is an overlap between our prediction for the unabsorbed source and/or the flared disk and the shaded region. Keep in mind, however, that the correction for absorption mentioned above aggravates the discrepancy between the model predictions and the observations (the latter being moved somewhat to the left in Fig. 8), but do not have effects on the conclusions on the scatter of [Ne~{\\sc ii}] luminosities.  It would be of little use here to speculate what other observational effects may be coming into play, however one thing that is worth noticing is that the measurements plotted in the Figure all come from the {\\it Spitzer} Space Telescope, while recent observations from the ground at higher resolution suggest often lower luminosities than the {\\it Spitzer} data \\citep{Herczeg2007,Najita2009}. This has been interpreted as possible pollution from molecular emission in the {\\it Spitzer} band, while more recent works (Najita et al., 2009) seem to lean toward extended emission, from undetected outflows, which would contribute to the flux in the {\\it Spitzer} aperture, but would be excluded by the narrow slits used for measurements from the ground.  Our results thus support a disk origin of the observed [Ne~{\\sc ii}] emission as a consequence of X-ray irradiation, at least for systems with moderate [Ne~{\\sc ii}]  luminosities and absence of jets. On the other hand, the additional parameters (X-ray spectral hardness, disk flaring) now found to influence the [Ne~{\\sc ii}] luminosity add complexity to the interpretation, and the usefulness of the [Ne~{\\sc ii}] line as a disk diagnostic depends on our ability to confine these other parameters as well. Clearly, [Ne~{\\sc ii}]  {\\it disk} diagnostics needs to avoid stars with known jets and outflows. Under the assumption of full dust and gas mixing, some information about disk flaring can be obtained by modeling the infrared spectral energy distribution \\citep{Chiang2001,Pascucci2003}, providing some important constraints on the range of modeled [Ne~{\\sc ii}]  luminosities. Constraining the shape of the ionizing spectrum is perhaps the most difficult task. The intrinsic spectrum can usually be modeled sufficiently well based on intermediate- or high-resolution X-ray spectroscopy, but the absorbing gas column density toward the disk may be different from that toward the observer. However, gas absorption matters the least in systems with weak or absent accretion, the possible presence of inner holes, and therefore they likely absence of strong winds emanating from the innermost disk regions. The [Ne~{\\sc ii}]  {\\it disk } diagnostic may thus be optimized for the study of transitional disks, most of which are also not known to drive outflows or jets.  Several of these have been well studied, including new ground-based observations of spectrally resolved  [Ne~{\\sc ii}] lines (Najita et al. 2009). In this context, the [Ne~{\\sc ii}] line may develop its full diagnostic power to study disk ionization and heating by ionizing radiation from the central star, processes that are pivotal for disk dispersal through photoevaporation \\citep{Alexander2006, Ercolano2008a, Ercolano2009, Owen2009}in the first place.  High resolution and high signal-to-noise line profile can help assessing the origin of the line, with profiles centred at the stellar radial velocity being synonymous with a disk origin. This is true for all inclinations except those very close to edge on, where outflows would also produce profiles centered on the stellar radial velocity. Blue shifted profiles are expected to be observed for system with outflows observed at non edge on inclinations (e.g. Alexander, 2008), and have been recently oberved by Pascucci et al. (2009) in the spectra of TW~Hya, Sz~73, T~Cha and CS~Cha. Other examples of high resolution spectra where the line has been detected include TW Hya (Herczeg et al. 2007; $\\lambda/\\Delta\\lambda \\sim$ 30000), GM Aur and AA Tau (Najita et al., 2009; $\\lambda/\\Delta\\lambda \\sim$ 80000). Here the observed lines are consistent with being centred at the stellar radial velocity, suggesting a circumstellar disk origin, although the poor signal-to-noise of some of these detections makes it hard to say with certainty. The lines are broad with a FWHM of respectively  $\\sim$ 21 (TW Hya), 70 (AA Tau) and 14 (GM Aur) km s$^{-1}$. The [Ne~{\\sc ii}] FWHM of GM Aur can be produced by a disk of normal flaring power, however the large FWHM of AA Tau requires the emission region to be dominated by very small radii and therefore implies a small degree of flaring. We finally conclude that high resolution observations of the [Ne~{\\sc ii}] line in YSOs, able to resolve its profile, are needed if this line is to be used to extract useful information on disk structure (e.g. flaring) and evolutionary stage by (e.g.) the detection of outflows and photoevaporative winds by comparison with line profile models like those of Alexander (2008) or those presented in this work.    \\begin{figure} \\centering \\includegraphics[width=9cm, trim = 25mm 125mm 15mm 25mm, clip]{./Figures/plot_LNeII_Lx_new.pdf} \\caption{Comparison between the prevision of our models, where a relationship between the L([Ne~{\\sc ii}]) and L$_X$ is expected, and the observed system with undetected outflows (from G\\\"udel 2008). } \\label{Fig:LNeIILX} \\end{figure}"
    },
    "0910/0910.2094_arXiv.txt": {
        "abstract": "{} {A nonlinear kinetic theory of cosmic ray (CR) acceleration in supernova remnants (SNRs) is employed to investigate the properties of SNR RX~J1713.7-3946.} {Observations of the nonthermal radio and X-ray emission spectra as well as the H.E.S.S. measurements of the very high energy \\gr emission are used to constrain the astronomical and CR acceleration parameters of the system. It is argued that \\rxj is a core collapse supernova (SN) of type II/Ib with a massive progenitor, has an age of $\\approx 1600$~yr and is at a distance of $\\approx 1$~kpc. It is in addition assumed that the CR injection/acceleration takes place uniformly across the shock surface for this kind of core collapse SNR.} {The theory gives a consistent description for all the existing observational data, including the nondetection of thermal X-rays and the spatial correlation of the X-ray and \\gr emission in the remnant. Specifically it is shown that an efficient production of nuclear CRs, leading to strong shock modification and a large downstream magnetic field strength $B_\\mathrm{d}\\approx 140$~$\\mu$G can reproduce in detail the observed synchrotron emission from radio to X-ray frequencies together with the \\gr spectral characteristics as observed by the H.E.S.S. telescopes.} {The calculations are consistent with \\rxj being an efficient source of nuclear cosmic rays.} ",
        "introduction": "RX~J1713-3946 is a shell-type supernova remnant (SNR) located in the Galactic plane that was discovered in X-rays with {\\it ROSAT} \\citep{pfe96}. Subsequent studies of this SNR with the {\\it ASCA} satellite by \\citet{koyama97} and  \\citet{sla99}, and later with {\\it XMM} by \\citet{cassam04}, have not positively identified any thermal X-ray emission and therefore tentatively concluded that the observable X-ray emission is entirely non-thermal.  The radio emission on the other hand is quite weak. Only part of the remnant shell could be detected in radio synchrotron emission up to now, with a poorly known spectral form \\citep{laz04}. For the spatially integrated radio flux \\citet{aha06} estimated about twice the value found by \\citet{laz04}. Most recently \\citet{acero09} derived a total flux value of $22 \\mathrm{Jy} < \\mathrm{S} < 26 \\mathrm{Jy}$ at 1.4~GHz, still about two times higher than estimated by \\citet{aha06}. As far as the estimate of the magnetic field strength is concerned, this higher estimate of the radio synchrotron flux increases the magnetic field estimate (see below).  RX~J1713-3946 was also detected in very high energy (VHE:$ >100$~GeV) \\grs with the {\\it CANGAROO} \\citep{mur00,enomoto02} and {\\it H.E.S.S.}  telescopes \\citep{aha04,aha06,aha07}. Especially the latter, very detailed observations show a clear shell structure at TeV energies which correlates well with the {\\it ASCA} contours.  In a first theoretical paper on this source \\citep{bv06} we have investigated the acceleration of CR electrons and protons in detail, using nonlinear kinetic theory \\citep{byk96,bv00}. Observations of the nonthermal radio and X-ray emission spectra as well as the H.E.S.S. measurements of the very high energy \\gr emission were used to constrain the astronomical and the particle acceleration parameters of the system. Under the assumption that \\rxj is the remnant of a core collapse supernova (SN) of type II/Ib with a massive progenitor, the theory gives indeed a consistent description for the existing data on the nonthermal emission. Specifically the VHE data from H.E.S.S. were shown to be best explained as hadronic \\gr emission, where the magnetic field amplification strongly depresses the inverse Compton and Bremsstrahlung fluxes. Subsequently \\citep{bv08}, we have analyzed the spatial correlation between the synchrotron and the VHE emission that has been used before as an argument for a leptonic origin of the VHE emission \\citep{Katz08,Plaga}. It was argued that correlated density and magnetic field strength variations lead rather naturally to a spatial correlation of the hadronic VHE emission with the synchrotron emission. Similar arguments were brought forward by \\citet{tua08}. Also the recent broad-band X-ray synchrotron measurements \\citep{uat07,ttu08} with the Suzaku instrument were compared to the theoretical synchrotron spectra and found to be quite consistent with a hadronic model. A purely leptonic model on the other hand \\citep[e.g.][]{Porter06}, where magnetic field amplification is not expected, was shown to be inconsistent with these observations. (\\citet{acero09} have also considered a leptonic origin of the \\gr emission as a viable possibility. This has to be compared with the results of the present paper.)  These theoretical considerations assumed that only over part of the shock surface injection of suprathermal nuclear ions occurs effectively, because over other parts the magnetic field is essentially parallel to the shock surface. In fairly regular SNRs like those of Type Ia explosions into a presumably uniform circumstellar environment, this is an important effect, as for instance the case of SN~1006 shows. For a blast wave propagating into the highly turbulent shell of a stellar wind bubble from a massive progenitor star on the other hand, the magnetic flux tubes in the shell are presumably so slender and irregular that energetic particles can cross effusively into regions where suprathermal injection is depressed, yet acceleration is possible. As a result shock modification by the pressure gradient of the accelerating particles becomes possible everywhere on the shock surface. This reasoning follows the discussion in \\citet{bpv09} with the conclusion that SNRs propagating into wind bubbles should have their forward shock modified everywhere. Technically this implies that the SNR becomes spherically symmetric and the correction factor $f_\\mathrm{re}<1$ \\citep{vbk03} on the spherically symmetric solution reaches unity. The ubiquitous shock modification has also implications for the thermal emission, because the thermal gas is then heated only in the subshock and the average gas density required for a given hadronic \\gr emission is reduced.  Contrary to the previous studies \\citep{bv06,bv08} of \\rxj we therefore adopt here the approximation of spherically symmetric CR injection/acceleration with $f_{\\mathrm{re}} \\approx 1$~ \\citep[see also][]{zirak09}. It is demonstrated that also in this approximation a consistent solution can also be found. In addition a rough estimate of the thermal X-ray emission from this object shall be given. Even though the error in this estimate might be quite large, the nominal result is consistent with the present non-detection of thermal X-rays \\citep[e.g.][]{acero09}.  Finally, the observed correlation of the X-ray and the VHE \\gr emission is discussed in a generalized form. ",
        "conclusions": "The assumption that CR injection/acceleration takes place uniformly across the whole SN shock surface is consistent with the existing data.  The swept-up mass is so small in this case that, in a rough approximation, the estimated flux of thermal X-rays at 1 keV is lower than the nonthermal flux, consistent with the nondetection of thermal X-ray emission until now.  It is concluded that the present observational knowledge of SNR \\rxj can be interpreted by a source which ultimately converts more than 35\\% of the mechanical explosion energy into nuclear CRs.  Also the observed high energy \\gr emission of SNR \\rxj turns out to be primarily of hadronic origin."
    },
    "0910/0910.1097_arXiv.txt": {
        "abstract": "The Balloon-borne Large-Aperture Submillimeter Telescope (BLAST) carried out a 250, 350 and 500\\,\\micron\\ survey of the galactic plane encompassing the Vela Molecular Ridge, with the primary goal of identifying the coldest dense cores possibly associated with the earliest stages of star formation. Here we present the results from observations of the Vela-D region, covering about $4\\,{\\rm deg}^2$, in which we find 141 BLAST cores.  We exploit existing data taken with the {\\it Spitzer\\/} MIPS, IRAC and SEST-SIMBA instruments to constrain their (single-temperature) spectral energy distributions, assuming a dust emissivity index $\\beta=2.0$.  This combination of data allows us to determine the temperature, luminosity and mass of each BLAST core, and also enables us to separate starless from proto-stellar sources. We also analyze the effects that the uncertainties on the derived physical parameters of the individual sources have on the overall physical properties of starless and proto-stellar cores, and we find that there appear to be a smooth transition from the pre- to the proto-stellar phase. In particular, for proto-stellar cores we find a correlation between the MIPS24 flux, associated with the central protostar, and the temperature of the dust envelope.  We also find that the core mass function of the Vela-D cores has a slope consistent with other similar (sub)millimeter surveys. ",
        "introduction": "Stars form in dense, dusty cores of molecular clouds, but little is known about their origin, their evolution, and their detailed physical properties. In particular, the relationship between the core mass function (CMF) and the stellar initial mass function (IMF) is poorly understood \\citep{mckee2007}.  One of the reasons for this lack of understanding, from the observational point of view, is the difficulty in selecting a statistically significant sample of truly {\\it pre-stellar} cores from an otherwise unremarkable collection of high column density features.  Pre-stellar cores represent a very early stage of the star formation (SF) process, before collapse results in the formation of a central protostar. The physical properties of these cores can reveal important clues about their nature; mass, spatial distributions, and lifetime are important diagnostics of the main physical processes leading to the formation of the cores from the parent molecular cloud. In addition, a comparison of the CMF to the IMF may help to understand what processes are responsible for further fragmentation of the cores, and thus the determination of stellar masses. Therefore, large samples of {\\it bona-fide} pre-stellar cores are important for comparison of observations with various SF models and scenarios.   Two major technological advances have revolutionized the study of the mass distribution in molecular clouds and the characterization of the earliest phases of SF: the Infrared Astronomical Satellite ({\\it IRAS}\\,) satellite, and the advent of submillimeter bolometer arrays such as MAMBO, SIMBA, and SCUBA. (Sub)millimeter surveys have been vitally important for probing the nature of SF, and in particular have represented the only way of deriving statistically-significant samples of pre- and proto-stellar objects. However, these instruments have probed the Rayleigh-Jeans tail of the spectral energy distribution (SED) of these cold objects, far from its peak. Therefore, these surveys have been limited by their relative inability to measure the temperature (e.g., \\citealp{motte1998}), producing large uncertainties in the derived luminosities and masses.  Recent surveys with the MIPS instrument of the {\\it Spitzer Space Telescope} are able to constrain the temperatures of warmer objects \\citep{car05}, but the youngest and coldest objects are potentially not detected, even in the long-wavelength {\\it Spitzer} bands.  Other ground-based surveys conducted at wavelengths $\\la 450 \\,$\\micron\\ have been affected by low sensitivity due to atmospheric conditions  at short submillimeter wavelengths (e.g., \\citealp{kirk2005, wu2007}).  The Balloon-borne Large-Aperture Submillimeter Telescope, BLAST  \\citep{pascale2008}, carried out two long duration balloon  science flights. During the first (BLAST05), BLAST mapped several galactic star forming regions \\citep{chapin2008, truch2008}. During the second science flight in 2006, BLAST06 mapped about  50\\,deg$^2$ in the Vela Molecular Ridge (VMR) \\citep{netterfield2009}. Until the more extensive results from {\\it Herschel} become available, BLAST is unique in its ability to detect and characterize cold dust emission from both starless and proto-stellar sources, constraining the temperatures of objects with $T \\la 25$\\,K using its three-band photometry (250, 350, and 500\\,\\micron) near the peak of the cold core SED.  The VMR is a giant molecular cloud complex within the galactic plane, in the area $260 \\degr \\la l \\la 272 \\degr$ and $-2 \\degr \\la b \\la 3 \\degr$, hence located outside the solar circle. It was first studied in detail by \\citet{mur:may}, using millimeter low-resolution observations in the CO(1--0) transition. They subdivided the VMR into four main regions, named A, B, C, and D.\\@ \\citet{Lis92} further discussed the issue of distance, finding that clouds A, C, and D are at $700 \\pm 200$ pc, whereas B seems to be at $\\sim 2000$ pc. In this work we have selected the Vela-D region because of the large amount of ancillary data available, both in the continuum, particularly in the mid- (MIR) and far-infrared (FIR), and also in terms of spectral line observations.  \\citet{Lis92} and \\citet{lor93} identified a number of VMR sites that host low- and intermediate-mass (proto)stars whose infrared SEDs are consistent with those of Class I sources. Later, \\citet{massi99, massi00, massi03} analyzed NIR images % of a sample of those sites, finding that the {\\it IRAS} sources with $L_{\\rm bol} > 10^{3} \\,$L$_{\\sun}$ are associated with small embedded young stellar clusters. Their IMF appears consistent with that of field stars and their age is of the order of $10^{6}$ yrs \\citep{massi06}. Remarkably, cloud D exhibits a lack of massive (O-type) stars with respect to what is expected from a standard IMF \\citep{massi06}.  \\begin{figure} \\centering \\includegraphics[width=\\linewidth]{TMmap1.eps} \\caption{ The gray-scale image shows the BLAST 250\\,\\micron\\ map of Vela-D, with galactic coordinates in degrees. Superimposed are the locations of both starless (open circles) and proto-stellar (crosses) cores (see \\S\\ref{sec:phot} and \\S\\ref{sec:comparison}). The size of both crosses and circles indicates the mass range of each core  and color-coding indicates the core temperature (see legend). See \\S\\ref{sec:masslum} for a discussion of the cores' physical parameters. } \\label{fig:blast250} \\end{figure}    New millimeter observations, both spectral line \\citep{elia2007}  and continuum \\citep{massi07}, have shown that all the {\\it IRAS} sources associated with embedded young clusters are also associated with one or more dense molecular cores, confirming the inferred young age of the clusters. On-going SF activity in Vela-D  is also confirmed by the presence of collimated jets originating from objects embedded in some of the clusters \\citep{lore02, delu07}. A first census of young stellar objects (YSOs) throughout the area of cloud D mapped by \\citet{massi07} and \\citet{elia2007}, was carried out by \\citet{gianni07} using {\\it Spitzer} MIPS observations at 24 and $70\\,$\\micron\\ (hereafter indicated as MIPS24 and MIPS70, respectively). Most of the millimeter cores identified by \\citet{massi07} were found to be associated with red, cold objects. Cloud~D appears therefore as an active star-forming region with both distributed and clustered SF in progress, hosting a large sample of young objects in different evolutionary stages.  In this paper we use both BLAST and archival data to determine the  SEDs, and thus physical parameters, of each source detected by BLAST. In \\S\\ref{sec:obs} we describe the BLAST observations and also the archival data that have been used in this work.  In \\S\\ref{sec:phot} we describe the identification of millimeter,  MIR and FIR counterparts of the BLAST cores, and also describe the multi-band photometry. This photometry is then used to construct the SED of each source in \\S\\ref{sec:sed}, which allow us to derive the physical parameters for starless and proto-stellar cores, discussed in \\S\\ref{sec:comparison}. We finally draw our conclusions in \\S\\ref{sec:concl}. ",
        "conclusions": "\\label{sec:concl}  In this paper we have presented a detailed analysis of the dense cores in the Vela-D molecular cloud, utilizing the sensitive maps at 250, 350 and 500\\,\\micron\\ obtained by BLAST during its 2006 science flight from Antarctica. We have combined the multi-color BLAST photometry with previous MIR, FIR, and millimeter-continuum observations to determine the physical parameters of the population of cores detected by BLAST. Our sample includes a total of 141 objects, although only 89 cores are located within the common area covered by the MIPS24 and MIPS70  maps. We summarize the other results of this paper as follows:  1. The multi-band photometry, combined with an isothermal modified blackbody model, allows us to analyze the effects that the uncertainties on the derived physical parameters of the individual sources have on the overall properties of starless and proto-stellar cores.  2. In the area of Vela-D covered by the MIPS24, MIPS70 and SIMBA maps we find 34 cores, out of 89, having no MIPS24 or MIPS70 counterpart, with 28 of these cores not having been detected by SIMBA either. Therefore, these observations demonstrate the importance of observing the early cores at or near the SED peak, as made possible by BLAST.  3. The CMF determined from all (starless and proto-stellar) BLAST cores in Vela-D shows a power-law slope of $\\alpha = -2.3 \\pm 0.2$, consistent, within the errors, with other continuum (sub)millimeter surveys and in particular with the slope found by \\citet{netterfield2009} in Vela-C.  4. We find a moderate difference between the median mass of proto-stellar and starless cores. We also find that proto-stellar cores are somewhat warmer and more luminous than starless cores, and that there is a progressively higher fraction of proto-stellar cores with increasing core temperature.  This suggests, as expected, a luminosity and temperature evolution due to the appearance of an embedded protostar.  5.  In terms of the observational properties, there appear to be a smooth transition from the pre- to the proto-stellar phase, as suggested by  differences found within the two populations. In particular, for proto-stellar cores we find a correlation between the MIPS24 flux, associated with the central protostar, and the temperature of the dust envelope. Our results can thus provide guidelines for understanding the physical conditions where the transition between pre- and proto-stellar cores takes place."
    },
    "0910/0910.4398_arXiv.txt": {
        "abstract": "Supermassive stars, with masses $\\ga 10^6 M_\\odot$, are possible progenitors of supermassive black holes in galactic nuclei.  Because of their short nuclear burning timescales, such objects can be formed only when matter is able to accumulate at a rate exceeding $\\sim 1 M_\\odot$ yr$^{-1}$.  Here we revisit the structure and evolution of rotationally-stabilized supermassive stars, taking into account their continuous accumulation of mass and their thermal relaxation.  We show that the outer layers of supermassive stars are not thermally relaxed during much of the star's main sequence lifetime. As a result, they do not resemble $n=3$ polytropes, as assumed in previous literature, but rather consist of convective (polytropic) cores surrounded by convectively stable envelopes that contain most of the mass.  We compute the structures of these envelopes, in which the equation of state obeys $P/\\rho^{4/3} \\propto M^{2/3}(R)$, where $M(R)$ is the mass enclosed within radius $R$.  By matching the envelope solutions to convective cores, we calculate the core mass as a function of time.  We estimate the initial black hole masses formed as a result of core-collapse , and their subsequent growth via accretion from the bloated envelopes (``quasistars\") that result.  The seed black holes formed in this way could have typical masses in the range $\\sim 10^4-10^5 M_\\odot$, considerably larger than the remnants thought to be left by the demise of Pop III stars.  Supermassive black holes therefore could have been seeded during an epoch of rapid infall considerably later than the era of Pop III star formation. ",
        "introduction": "Supermassive stars were proposed by Hoyle \\& Fowler (1963a,b) as a means to meet the prodigious energy requirements of radio galaxies (Burbidge 1958), and slightly later as a model for quasars.  Hoyle \\& Fowler recognized that such objects could not persist for much more than a million years, if sustained by hydrogen burning, and suggested that they might collapse to form black holes once their nuclear fuel was exhausted.  The pulsational stability of supermassive stars is a crucial issue, because they are radiation pressure-dominated and therefore have an adiabatic index very close to $4/3$ (which yields neutral stability to radial pulsations for a Newtonian, self-gravitating body with no rotation).  Small general relativistic corrections have a destabilizing effect, preventing nonrotating stars more massive than a few $\\times 10^5 \\ M_\\odot$ from attaining a phase of stable hydrogen burning before collapsing (Iben 1963; Fowler 1964).  However, a dynamically insignificant level of rotation --- especially differential rotation --- can stabilize stars as massive as $\\sim 10^8 M_\\odot$ or more (Fowler 1966).  Additional stabilizing effects due to magnetic fields and turbulence were later considered by other authors (e.g., Bisnovatyi-Kogan, Zel'dovich \\& Novikov 1967;   Ozernoy \\& Usov 1971).  In this paper, we are interested in whether stable supermassive stars might be precursors of seed black holes that eventually grow to large enough mass to power quasars and populate the nuclei of present-day massive galaxies.  As we will argue in \\S~5, the holes created inside a supermassive star are likely to have masses of a few percent of the star's final mass.  Thus, to obtain seed black holes with masses $\\sim 10^4-10^5 M_\\odot$, comfortably larger than the seeds probably left behind by the collapse of Population III stars, we need to consider supermassive stars with masses $\\ga 10^6 M_\\odot$.  We will focus on this mass range in our analysis.\\footnote{Note, however, that the term {\\it supermassive star} is often used to refer to the mass range $M_*\\ga 5\\times 10^4  M_\\odot$ (Fuller, Woosley \\& Weaver 1986).}  Once the existence of a stable supermassive star is postulated, its structure would appear to be quite simple. In order for radiation pressure to support the star against gravity, the radiative flux seen by every element of mass in the star must equal the {\\it local} Eddington limit, $F_E =  GM(R)c/\\kappa r^2$, where $\\kappa$ is the local opacity and $M(R)$ is the mass enclosed within $R$.  The opacity inside a supermassive star is likely to be roughly uniform, with electron scattering dominant everywhere, but $M(R)$ is a monotonic function of radius.  Since the entire luminosity is produced by thermonuclear reactions within a small region of the core, the star must be highly convective with the fraction of the total luminosity carried by convection varying as $L_{\\rm conv}/L_{\\rm tot} = 1 - M(R)/M_*$, where $M_*$ is the total mass of the star.  If the convection is efficient, then a supermassive star should be described accurately by an $n=3$ ($\\gamma = 4/3$) polytrope, according to the Lane-Emden equation.  Such models, however, beg the question of how a supermassive star might realistically form.  Creating a $10^6 M_\\odot$ star requires the very rapid accumulation of gas. Since the thermonuclear timescale for a star burning hydrogen at the Eddington limit is $\\sim2$ Myr, independent of mass, and will be less if only part of the fuel is burned, the infall rate required to create a supermassive star of mass $M_*$ must exceed $0.5 (M_*/10^6 M_\\odot) \\ M_\\odot$ yr$^{-1}$.  This is orders of magnitude larger than the rate at which matter condenses to form normal stars within molecular clouds, or even the rate ($\\sim 10^{-3} M_\\odot$ yr$^{-1}$) at which matter came together to form Pop III stars in pregalactic dark matter halos.  The maximum rate at which matter can collect is given by $\\sim v^3/G = 0.2 (v/10 \\ {\\rm km \\ s}^{-1})^3 M_\\odot$ yr$^{-1}$, under the assumption that the gas is self-gravitating and collapses at its free-fall speed $v$.  In effect, this means that supermassive stars can only form in systems with virial temperatures exceeding $10^4$ K.  There is considerable controversy over whether such high rates of inflow can occur without most of the gas fragmenting and forming stars before reaching the center.  Despite early suggestions that fragmentation is avoidable only if the gas temperature remains close to the virial temperature (Bromm \\& Loeb 2003; Begelman, Volonteri \\& Rees 2006) --- and thus only in systems lacking both metals and molecular hydrogen ---  recent simulations suggest that fragmentation may be suppressed even when the gas is much colder (Wise, Turk \\& Abel 2008; Levine et al. 2008; Regan \\& Haehnelt 2009), possibly due to the continuous generation of supersonic turbulence (Begelman \\& Shlosman 2009).  Conditions of rapid infall are ideal for producing the high levels of entropy required inside supermassive stars.  Indeed, simple estimates show that the initial entropy of the gas joining the star can be much larger than the value required for equilibrium.  We argue that the nature of radiation-pressure support provides a mechanism for automatically regulating the entropy of the gas joining the supermassive star, so that it is neither too large nor too small. This is due to the existence of a ``trapping radius\" within the infalling gas, outside of which radiative diffusion can release excess entropy.  Gas with too much entropy will be forced to expand until it gives up the excess. However, this does {\\it not} imply that supermassive stars should have uniform specific entropy and therefore resemble polytropes.  In this paper we study the evolution and fate of supermassive under the assumption that they grow by rapid, continuous infall. In \\S\\S~2 and 3, we argue that gas joins the star with increasing specific entropy as a function of time.  Since the Kelvin-Helmholtz timescale is longer than the age of the star for early times and high infall rates, much of this this entropy stratification is preserved over the life of the star.  We therefore conclude that supermassive stars are not necessarily well-represented by $n=3$ polytropes, but rather can have a more complex structure with a convective (polytropic) core surrounded by a convectively stable envelope that contains most of the mass.  The entropy in the envelope roughly satisfies $P/\\rho^{4/3} \\propto M^{2/3}(r)$ (we term this entropy law ``hylotropic\").  Hydrogen burning in the core starts when the star's mass and entropy are both relatively low, and adjusts to sustain the star through its more massive stages (\\S~4).  In \\S~5, we discuss the formation of the seed black hole after the exhaustion of core hydrogen, and its subsequent growth inside the remnant of the supermassive star.  Because of the likely importance of rotation in the collapsing core, only a small fraction of the supermassive star collapses to a black hole initially.  The energy liberated during the formation of the black hole inflates the remainder of the star into a bloated object that resembles a red giant, which we have previously termed a ``quasistar\" (Begelman et al.~2006; Begelman, Rossi \\& Armitage 2008).  The black hole continues to grow by accretion from the quasistar envelope, until the photospheric temperature drops to the point where the quasistar undergoes an ``opacity crisis\" and disperses under the influence of radiation pressure.  We summarize our results in \\S~6 and comment on the possible contribution of supermassive stars to the ionizing radiation field at high redshifts. ",
        "conclusions": "We have revisited the theory of supermassive stars, taking into account the fact that their short nuclear burning timescales ($\\leq 4$ Myr) imply that the matter forming them must accumulate rapidly.  If the accumulation of matter is rapid enough (typically, more than a few solar masses per year), then the star never has time to become thermally relaxed and fully convective. Instead, its structure consists of a convective core containing a minority of the mass and occupying a small fraction of the stellar volume, surrounded by a convectively stable envelope with the specific entropy increasing outward as enclosed mass to the 2/3 power. We have calculated the structures of these envelopes --- for which we propose the term ``hylotropes\" (following a suggestion by A.~Accardi and G.~Lodato) --- and the conditions for matching them to the convective cores, and have thus devised a theory for the evolution of partially convective supermassive stars.  Partially convective supermassive stars are strikingly different from their fully convective counterparts (described by $n=3$ polytropes) in a number of respects.  They do not follow the same mass-radius relation as fully convective stars.  Instead of the scaling $R \\propto M^{1/2}$, which applies for a fully convective supermassive star with a central temperature regulated by thermonuclear reactions, our partially convective models have a radius proportional to the (assumed constant) accretion rate.  Physically, the stellar radius is comparable to the ``trapping radius\" within which radiation is unable to escape inward advection by the accumulating gas (Begelman 1978, 1979). Also, their nuclear-burning (``main sequence\") lifetimes can be shorter because the fuel in the convectively stable envelope is never mixed into the core.  Hydrogen-burning supermassive stars in the mass range considered here ($\\ga 10^6 M_\\odot$) can exist only if stabilized by rotation or some other ``stiff\" form of energy (such as magnetic fields or turbulence) (Fowler 1964, 1966).  The minimal level of rotation required to cancel the general-relativistic instability is not dynamically significant.  However, the angular momentum may well be high enough to inhibit the formation of the black hole, once the core runs out of hydrogen and collapses.  We have argued that the initial mass of the seed black hole, if limited by the energy released in forming it, may be as small as a few hundred solar masses --- i.e., only $10^{-4}$ of the total stellar mass.  However, there is likely to be a second phase of rapid black hole growth once the inflated stellar envelope reaches equilibrium as a red giant-like quasistar.  We therefore estimate that supermassive stars should leave behind black holes with typical masses of a few percent that of the star, i.e., roughly $\\sim 10^4 - 10^5 M_\\odot$.  These estimates are admittedly quite uncertain, mainly because we do not understand how energy liberated by black hole accretion couples to the remainder of the star.  We cannot rule out the possibility that most of the energy punches through the star in a pair of jets (in a fashion similar to the collapsar model  of gamma-ray bursts [Woosley 1993]), in which case both the seed and final black hole masses could be much larger.   We note that while our estimate of the initial black hole mass is qualitatively similar to values (few tens of $M_\\odot$) predicted by Begelman et al.~(2006), our new value for the predicted envelope mass at the time of black hole formation is 2--3 orders of magnitude larger.  The reason for this discrepancy is that Begelman et al.~(2006) failed to account for the rapid growth of the convective core due to heating by efficient thermonuclear reactions in the CNO cycle.  (Effectively, we considered pure Pop III abundances, for which thermonuclear energy release is too weak to stall the formation of the seed hole.)  Supermassive stars can be strong sources of far-UV radiation, provided that their spectra are not severely degraded by reprocessing at radii beyond $R_{\\rm tr}$.  Taking the photospheric radius to be $0.6 R_{\\rm tr}$, we find an effective temperature $T_{\\rm eff} = 1.2 \\times 10^5 m_{*,6}^{1/4} \\dot m_*^{-1/2}$ K for a partially convective star of mass $M_*$, corresponding to $1.3 \\times 10^5 \\dot m_*^{-3/8} T_{c,8}^{-1/4}$ K at the end of hydrogen-burning.  However, the infalling gas joining a supermassive star is already very optically thick at $R_{\\rm tr}$ ($\\tau \\sim c/v$, where $v$ is the infall speed), implying that most of the radiation escapes from further out and is likely to be softer. On the other hand, the opacity is strongly scattering-dominated, so the color temperature could be several times higher than the effective temperature.  The hardest radiation is likely to escape from the polar regions, where rotational effects decrease the density.  Assuming a blackbody spectrum with the effective temperature estimated above and $\\dot m_* = T_{c,8} = 1$, we find that most of the radiation is capable of ionizing hydrogen ($E > 13.6$ eV), with an ionizing luminosity $L_{\\rm ion}\\approx 2\\times 10^{44}$ erg s$^{-1}$ and photon flux ${\\cal N}_{\\rm ion} \\approx 4 \\times 10^{54}$ s$^{-1}$. The star radiates into the Lyman--Werner band ($11.2 < E < 13.6$ eV), capable of dissociating H$_2$, with $L_{\\rm diss}\\approx 5\\times 10^{42}$ erg s$^{-1}$ and ${\\cal N}_{\\rm diss} \\approx 3 \\times 10^{53}$ s$^{-1}$.  To produce an interesting mean intensity of dissociating radiation, $J_\\nu > 10^{-22}$ erg cm$^{-2}$ s$^{-1}$ Hz$^{-1}$ sr$^{-1}$ between $z\\sim 12$ and $z\\sim 6$, would require (very roughly) a comoving density of supermassive stars of at least $\\sim 10^{-4}$  Mpc$^{-3}$.  Given the short lifetimes of these objects, it is not clear how plausible this is.  We will investigate cosmological scenarios for the formation of supermassive stars in a future publication."
    },
    "0910/0910.1604_arXiv.txt": {
        "abstract": "We present a near-infrared spectroscopic study of HD 114762B, the latest-type metal-poor companion discovered to date and the only ultracool subdwarf with a known metallicity, inferred from the primary star to be [Fe/H] = --0.7.  We obtained a medium-resolution (R $\\sim$ 3800) Keck/OSIRIS 1.18-1.40 $\\mu$m spectrum and a low-resolution (R $\\sim$ 150) IRTF/SpeX 0.8-2.4 $\\mu$m spectrum of HD 114762B to test atmospheric and evolutionary models for the first time in this mass-metallicity regime.  HD 114762B exhibits spectral features common to both late-type dwarfs and subdwarfs, and we assign it a spectral type of d/sdM9 $\\pm$ 1.  We use a Monte Carlo technique to fit PHOENIX/$GAIA$ synthetic spectra to the observations, accounting for the coarsely-gridded nature of the models.  Fits to the entire OSIRIS $J$-band and to the metal-sensitive $J$-band atomic absorption features (\\ion{Fe}{1}, \\ion{K}{1}, and \\ion{Al}{1} lines) yield model parameters that are most consistent with the metallicity of the primary star and the high surface gravity expected of old late-type objects.  The effective temperatures and radii inferred from the model atmosphere fitting broadly agree with those predicted by the evolutionary models of Chabrier \\& Baraffe, and the model color-absolute magnitude relations accurately predict the metallicity of HD 114762B.  We conclude that current low-mass, mildly metal-poor atmospheric and evolutionary models are mutually consistent for spectral fits to medium-resolution $J$-band spectra of HD 114762B, but are inconsistent for fits to low-resolution near-infrared spectra of mild subdwarfs.   Finally, we develop a technique for estimating distances to ultracool subdwarfs based on a single near-infrared spectrum.  We show that this ``spectroscopic parallax'' method enables distance estimates accurate to  $\\lesssim$ 10\\% of parallactic distances for ultracool subdwarfs near the hydrogen burning minimum mass. ",
        "introduction": "A complete understanding of low-mass stars and brown dwarfs is an important goal of astrophysics.  Low-mass stars (0.08 $M_{\\odot}$ $\\lesssim$ $M$ $\\lesssim$ 0.5 $M_{\\odot}$)  are the most numerous objects in our galaxy (\\citealt{Lada:2006p14060}) and make up nearly half of its total mass. Brown dwarfs are objects that form like stars but have insufficient mass to sustain core nuclear reactions ($M$ $\\lesssim$ 0.08 $M_{\\odot}$) and represent the link between low-mass stars and extra-solar giant planets.  Objects with spectral types of M7 and later are known as ``ultracool'' objects and encompass the lowest-mass stars and brown dwarfs.  Our understanding of these objects has dramatically improved over the last decade, but there remain several fundamental areas of parameter space that are relatively unexplored, one of which is metallicity.   Sub-solar metallicity low-mass stars and brown dwarfs are referred to as ultracool subdwarfs.  Like their solar-metallicity counterparts, ultracool subdwarfs have spectral energy distributions that peak in the near-infrared and are dominated by overlapping molecular absorption bands.  Their lower metallicities, however, dramatically alter their spectral properties compared to solar-metallicity objects and can ultimately provide insight into the influence of metallicity on the chemistry, condensate cloud formation, and temperature/pressure profiles in the atmospheres of ultracool objects.  Prominent spectral changes resulting from a reduced metallicity include suppressed $H$- and $K$-band fluxes (i.e. bluer near-infrared colors) due to increased collision-induced absorption by H$_2$ (CIA H$_2$; \\citealt{Linsky:1969p3947}; \\citealt{Borysow:1997p3945}), a suppression of metal-oxide bands, and an enhancement of metal hydride bands (\\citealt{Bessell:1982p13345}).  Ultracool subdwarfs exhibit large space motions consistent with thick disk or halo kinematics (\\citealt{Burgasser:2007p575}), indicating that these objects are probably quite old (see, e.g., \\citealt{Helmi:2008p17636}).  \\begin{deluxetable}{lcc} \\tabletypesize{\\scriptsize} \\tablewidth{0pt} \\tablecolumns{3} \\tablecaption{Age and Metallicity of HD 114762A: \\\\ Literature Search \\label{metallicity} (1995 -- 2008)} \\tablehead{ \\colhead{Reference}    &  \\colhead{ Age (Gyr) }   &  \\colhead{ [Fe/H] } }  \\startdata  \\citet{Haywood:2008p11758}  & 12.4 (2.8)   &  \\nodata     \\\\ \\citet{Holmberg:2007p11970} & 10.6 (2.2)  &  --0.76  \\\\ \\citet{Gonzalez:2007p11922}   & \\nodata   & --0.657 (0.051) \\\\ \\citet{Zhang:2006p197}    & 14.1 &      --0.80 (0.10)  \\\\ \\citet{Reddy:2006p11923}  &  11.3 (3.0)   &  --0.71  \\\\ \\citet{Saffe:2005p183}   &  11.8 (3.6)   &   \\nodata  \\\\ \\citet{Valenti:2005p11833}   &   7.7 (2.3)   &  --0.65 (0.02)  \\\\ \\citet{Santos:2004p5955}      & \\nodata  &  --0.70 (0.04)    \\\\ \\citet{Gratton:2003p6816}   &  \\nodata & --0.79$_\\mathrm{Fe \\ I} $ \\\\ \\citet{Laws:2003p12112} & 14 (2)  &  \\nodata \\\\ \\cite{Heiter:2003p11754}  &  \\nodata  &  --0.78 (0.06) \\\\ \\cite{Sadakane:2002p12239} & \\nodata & --0.74 (0.07) \\\\ \\citet{Fulbright:2000p11924}  &  \\nodata  &  --0.7 (0.08) \\\\ \\citet{Lachaume:1999p11806} & 11 (1)  & \\nodata \\\\ \\citet{Thevenin:1999p11819} & \\nodata & --0.72$_\\mathrm{LTE}$; --0.56$_\\mathrm{NLTE}$ \\\\ \\citet{Clementini:1999p11805}  &  \\nodata  &  --0.66 (0.07) \\\\ \\citet{Gonzalez:1998p5871}  &  14 (2) & --0.6 (0.06)    \\\\ \\citet{Fuhrmann:1998p12487} &   \\nodata   &  --0.71 ($\\sim$0.05) \\\\ \\citet{Ng:1998p5790}  &  10.8 (0.8)  & \\nodata  \\\\ \\citet{Henry:1997p5745}  &  5   &\\nodata  \\\\ \\cline{1-3} \\\\ \\textbf{Adopted Values\\tablenotemark{a}}  &  \\textbf{11 (3)}  & \\textbf{--0.71 (0.07)} \\\\  \\enddata \\tablenotetext{a}{The adopted age and metallicity represent the mean and rms value of the listed quantities.}  \\end{deluxetable}  The number of spectroscopically-confirmed ultracool subdwarfs has rapidly increased from the first identification over a decade ago (LHS 377, with a spectral type of sdM7: \\citealt{Monet:1992p2839}; \\citealt{Gizis:1997p80}) to the $\\sim$ 45 currently known, the vast majority of which have been found in the past five years (see Table 7 of \\citealt{Burgasser:2007p575}; \\citealt{Burgasser:2008p2475}).  Many discoveries have been made through searches of large near-infrared and proper motion surveys including the Deep Near Infrared Survey of the Southern Sky (\\citealt{Epchtein:1997p4581}), the Two Micron All Sky Survey (2MASS, \\citealt{Skrutskie:2006p589}), and the Digitized Sky Survey (\\citealt{Lepine:2005p70}).  For example, \\citet{Lepine:2008p2934} recently identified  23 new M-type ultracool subdwarfs through template matching to Sloan Digital Sky Survey optical spectra.  Many more field ultracool subdwarfs are expected to be revealed with the next generation of deep, multi-epoch all-sky surveys such as the Panoramic Survey Telescope \\& Rapid Response System (\\citealt{Kaiser:2002p14879}) and the Large Synoptic Survey Telescope (\\citealt{Tyson:2002p14903}).   The absolute metallicities of ultracool subdwarfs are currently unknown.  A commonly-used technique for estimating the metallicities of late-type objects is to fit synthetic spectra to their red-optical spectra and to adopt the metallicity of the best-fitting model (\\citealt{Schweitzer:1999p563}; \\citealt{Lepine:2004p559}; \\citealt{Burgasser:2007p575}).   Evolutionary models have also been used to estimate the metallicities of ultracool subdwarfs in color-color  and color-magnitude diagrams (\\citealt{Scholz:2004p580}; \\citealt{Burgasser:2008p2475}; \\citealt{Dahn:2008p14767}; \\citealt{Schilbach:2009p14779}).  It is unclear how reliable these methods are, however, because models in this low-temperature, low-metallicity regime have not been tested.  There is thus a growing need for the discovery of ultracool subdwarfs with independently derived metallicities.   Benchmark objects whose fundamental parameters can be independently determined enable valuable tests of atmospheric and evolutionary models.  Several solar metallicity benchmark systems have already been used for such purposes.  One of the best examples is the HD 130948ABC system (\\citealt{Potter:2002p5031}), which contains two L4 brown dwarfs in a hierarchical triple system with a G2V primary.  In this case the age and metallicity of the primary star along with the luminosities and dynamical masses of the brown dwarfs have been directly measured and used to test brown dwarf evolutionary models (\\citealt{Dupuy:2009p15627}).  At sub-solar metallicities, however, there have not been any benchmark studies using very low-mass objects.  One possibility would be to use the latest-type objects in globular clusters, which comprise coeval stellar populations with uniform chemical compositions.  Their great distances, however, currently prevent spectroscopic studies of the lowest-mass members for which even photometric detections have proven difficult to obtain (\\citealt{Richer:2002p14764}; \\citealt{Richer:2006p14724}; \\citealt{Richer:2008p14712}).  Another approach would be to use nearby ultracool subdwarfs in binary systems as sub-solar metallicity benchmarks.  Companions to field subdwarfs could also act as metallicity calibrators for the field ultracool subdwarf population.   This technique relies on the conservative assumption that the components of binary systems have the same metallicity.  Observations of binaries in the field (\\citealt{Desidera:2004p109}) and in the Hyades (\\citealt{Paulson:2003p3765}) have revealed that the components had differential [Fe/H] abundances of $\\lesssim$ 0.02-0.04 dex, implying that the assumption of identical metallicity is indeed reasonable.   Unfortunately, most of the known ultracool subdwarfs are isolated objects and searches for low-mass secondaries to metal-poor stars have yielded surprisingly few bona fide companions (\\citealt{ZapateroOsorio:2004p5007}; \\citealt{Chaname:2004p4964}; \\citealt{Riaz:2008p4770}).      HD 114762B is currently the latest-type object known that is gravitationally bound to a metal-poor star (sdF9; HD 114762A).   It was discovered and confirmed as a proper motion companion in a Lick Observatory adaptive optics search for stellar and substellar companions to exoplanet host stars (\\citealt{Patience:2002p186}) and was later rediscovered in a similar study using CFHT AO (\\citealt{Chauvin:2006p178}).  HD 114762B is located 3$\\farcs$3 (128 AU at 38.7 pc; \\citealt{vanLeeuwen:2007p12454}) from its primary, which also harbors a high-mass planet, brown dwarf, or low-mass star, depending on the system's inclination ($M_\\mathrm{P}$sin$i$ = 11.68 $\\pm$ 0.96 $M_\\mathrm{Jup}$; \\citealt{Latham:1989p5199}; \\citealt{Cochran:1991p5246}; \\citealt{Hale:1995p6074}; \\citealt{Butler:2006p3743}).  The presence of a massive exoplanet candidate has led to numerous studies of the primary's metallicity ([Fe/H]) and age, which have been well constrained to be --0.71 $\\pm$ 0.07 dex and 11 $\\pm$ 3 Gyr, respectively\\footnote{The quoted errors are rms values based on a literature search rather than formal uncertainties.}  (Table \\ref{metallicity}).  Age estimates in the literature for HD 114762A are based on theoretical isochrones and chromospheric activity levels.  This system may also harbor a debris disk based on optical polarimetry and ISO far-IR photometry (\\citealt{Saffe:2004p187}).  A Keck adaptive optics/NIRSPEC $J$-band spectrum of HD114762B was presented by \\citet{Patience:2002p186}.  They found inconsistent spectral features that appeared similar to a late-M/early L spectrum based on the strength of the 1.4 $\\mu$m steam band, but also resembled a mid-M spectrum based on the strengths of the 1.189 $\\mu$m \\ion{Fe}{1} and 1.313 $\\mu$m \\ion{Al}{1} neutral metal lines and on a comparison to $J$-band spectra of other M dwarfs.  However, the dearth of ultracool subdwarfs known at the time of discovery prevented a comparative analysis with other metal-poor objects, and the nature of their observations made the analysis difficult because slit-based spectroscopy with adaptive optics does not preserve the continuum shape  (\\citealt{Goto:2002p495}; \\citealt{Goto:2003p13421}; \\citealt{McElwain:2007p13510}).   Here we present a near-infrared spectroscopic study of HD 114762B.  Our goals are threefold: ($1$) to better characterize this object through a comparison to known subdwarfs, ($2$) to test atmospheric models by comparing the results of the best-fitting models to the known metallicity of this object, and (3) to study the consistency of the predictions from atmospheric and evolutionary models.  For our tests we use the $GAIA$ grid of synthetic spectra (\\citealt{Brott:2005p301}, based on the PHOENIX atmospheric model code) and the \\citet[CB97]{Chabrier:1997p2767}, \\citet[BCAH97]{Baraffe:1997p582}, and \\citet[BCAH98]{Baraffe:1998p160}  low-metallicity evolutionary models (all based on the same input physics).  These models are widely used throughout the literature and are therefore important to validate.  In $\\S$2 we present our observations and the methods for extracting the spectra from the reduced data.  In $\\S$3 we describe the results of the $\\chi^2$ fitting of the models to the data and the way in which we derive the fundamental parameters from the atmospheric and evolutionary models.  Finally, we summarize our results and discuss the implications of our work in $\\S$4. ",
        "conclusions": "We have presented a near-infrared spectroscopic analysis of HD 114762B, the first benchmark ultracool subdwarf companion.   The metallicity of HD 114762B is inferred from the primary star, which has a mean metallicity of [Fe/H] = --0.71 based on independent estimates from the literature.  The spectral characteristics of HD 114762B include suppressed $H$ and $K$-band continuum as well as deeper 1.4 and 1.9 $\\mu$m H$_2$O bands compared to solar-metallicity M dwarfs.  The SpeX 1.00--2.35 $\\mu$m spectrum does not, however, exhibit the extreme CIA H$_2$ which is characteristic of the ``sd'' class of ultracool subdwarfs.  In this respect, HD 114762B is more similar to mild subdwarfs, which are believed to have an intermediate metallicity between dwarfs and subdwarfs.  We assign it a spectral type of d/sdM9 $\\pm$ 1 based on its resemblance to 2MASS 0041+35 (d/sdM9) from 1.00 to 2.35 $\\mu$m.  A spectral type later than sdM7.5 can also be inferred  from the strength of the \\ion{K}{1} lines and the strength of the 1.4 $\\mu$m H$_2$O band compared to M-type subdwarf spectra.  The measured luminosity places HD 114762B just above the hydrogen-burning minimum mass based on the evolutionary models of CB97.  We use this unique  benchmark object to test low-mass metal-poor atmospheric and evolutionary models.  The best-fitting PHOENIX/$GAIA$ synthetic spectrum to our OSIRIS $J$-band data has physical parameters that are consistent with both the metallicity of the primary star and the high surface gravity of late-type objects.  Fits to the SpeX 1.00--2.35 $\\mu$m spectrum, however, yield physical parameters that are inconsistent with these values.  We also fit the models to low-resolution spectra of other late-type mild ultracool subdwarfs with similar spectral types to HD 114762B.  The best-fitting models' surface gravity and metallicity are mutually inconsistent among the four objects and were likewise inconsistent for fits to different regions of the same spectrum.  Fits to HD 114762B, however,  yield consistent effective temperatures between 2500-2800 K.  We conclude that the metallicities and surface gravities derived from fitting atmospheric models to low-resolution near-infrared spectra of ultracool subdwarfs are not trustworthy, but that the effective temperatures are probably more reliable.  Problems with atmospheric models at low masses and low metallicities have already been noted in the literature; only moderately good matches have resulted from fits to the red-optical spectral regions of ultracool subdwarfs, even when applied over relatively short wavelength ranges (\\citealt{Schweitzer:1999p563}; \\citealt{Lepine:2004p559}; \\citealt{Burgasser:2007p575}).   These same studies held the  model surface gravity fixed at log $g$ = 5.0 or 5.5 while only allowing the temperature and metallicity to vary.  The models should, however, be able to predict all three parameters with no \\emph{a priori} assumption about any particular one.  In addition, the models should be able to predict roughly the same parameters for objects of the same spectral type as well as for different spectral regions of the same object.  We have shown that neither appears to be true for the ultracool subdwarfs we studied, making the discovery of benchmark metallicity calibrators like HD 114762B ever more important.  Our results suggest that the best method to determine the metallicity of field ultracool subdwarfs may be to use the medium-resolution $J$-band region.  The 1.18-1.35 $\\mu$m range encompasses metal-sensitive atomic absorption lines as well as the metal-sensitive 1.4 $\\mu$m steam band.  An  interesting line that may be important for very metal-poor ultracool subdwarfs is the \\ion{Fe}{1} 1.887 $\\mu$m feature, which the $GAIA$ models predict to be essentially independent of temperature and gravity, but which is strongly dependent on metallicity for [M/H] $\\lesssim$ --1.0.  We found that fitting the individual absorption lines rather than the entire 1.18-1.35 $\\mu$m spectrum of HD 114762B produced temperatures and gravities that were more consistent with the values derived from the CB97 evolutionary models.  Nevertheless, fitting the entire $J$-band spectrum and the individual absorption features yielded metallicities consistent with the primary star and surface gravities consistent with empirical values of late-type objects.  The location of HD 114762B in the BCAH97/98 M$_{K_\\mathrm{S}}$ vs $J$ -- $K_\\mathrm{S}$ color-absolute magnitude diagram is consistent with the metallicity of the primary star.   The metallicity and surface gravity inferred from $GAIA$ fits to low-resolution near-infrared spectra are not trustworthy, but the results from the color-magnitude diagram suggest that the overall shape of the synthetic spectra are in fact reliable, at least for the mass ($\\sim$ 0.088 $M_{\\odot}$) and metallicity ([Fe/H] = --0.71) of HD 114762B.  Additionally, we calculated the bolometric luminosities and bolometric corrections for four known ultracool subdwarfs (LHS 377, LSR 2036+5059, SSSPM 1013--1356, and 2MASS 1626+3945) with parallaxes and whose near-infrared spectra were available in the SpeX Prism Spectral Library.  Ultracool subdwarfs have higher luminosities and smaller bolometric corrections than do ultracool dwarfs in the same spectral subclass.  While unresolved binarity is a possibility, a hotter effective temperature scale and/or a larger radius is a more likely explanation for the overluminosity.  Finally, we have also developed a  technique to estimate the distances to ultracool subdwarfs based on the available models and a single near-infrared spectrum.  Using the effective temperature from the best-fitting synthetic spectrum and an assumed age ($\\sim$ 10 Gyr for objects with thick disk or halo kinematics), a radius can be inferred from evolutionary models.  The radius can then  be used along with the model scaling factor $\\overline{C}_k$ (equal to $R^2/d^2$) to derive a distance estimate to the object.  We applied this technique to five ultracool subdwarfs with trigonometric distances and obtained distance estimates accurate to $<$ 10\\% of the parallactic distances, or about 2 times better than estimates based on photometry alone.  This technique is particularly useful for ultracool subdwarfs, whose ages can be constrained by their kinematics.   It will be worthwhile to study HD 114762B at other wavelengths, including obtaining thermal infrared photometry ($\\lambda$ $>$ 3 $\\mu$m) to better constrain the long-wavelength end of the model-fitting, and, if possible, optical spectra, which would allow for an independent spectral classification of this object.  The dominant error in our luminosity determination originates from the errors in the photometry used to flux calibrate the spectra.  Better near-infrared photometry would therefore enable more a precise luminosity estimate and, by extension, more precise predictions from the evolutionary models."
    },
    "0910/0910.0214_arXiv.txt": {
        "abstract": "We report the detection of a strong, organised magnetic field in the O9IV star HD~57682, using spectropolarimetric observations obtained with ESPaDOnS at the 3.6-m Canada-France-Hawaii Telescope within the context of the Magnetism in Massive Stars (MiMeS) Large Program. From the fitting of our spectra using NLTE model atmospheres we determined that HD~57682 is a $17^{+19}_{-9}$~M$_{\\odot}$ star with a radius of $7.0^{+2.4}_{-1.8}$~R$_\\odot$, and a relatively low mass-loss rate of $1.4^{+3.1}_{-0.95}\\times10^{-9}$~M$_{\\odot}$\\,yr$^{-1}$. The photospheric absorption lines are narrow, and we use the Fourier transform technique to infer $v\\sin i=15\\pm3$~km\\,s$^{-1}$. This $v\\sin i$ implies a maximum rotational period of 31.5 d, a value qualitatively consistent with the observed variability of the optical absorption and emission lines, as well as the Stokes~$V$ profiles and longitudinal field. Using a Bayesian analysis of the velocity-resolved Stokes~$V$ profiles to infer the magnetic field characteristics, we tentatively derive a dipole field strength of $1680^{+134}_{-356}$~G. The derived field strength and wind characteristics imply a wind that is strongly confined by the magnetic field. ",
        "introduction": "The phenomenon of magnetism in hot, massive OB stars is not well studied. To date, repeated detections of circular polarization within line profiles have firmly established the presence of magnetic fields in three O-type stars ($\\theta^1$~Ori~C, HD~191612, and $\\zeta$~Ori~A; Donati et al. 2002, 2006; Bouret et al. 2008) and a handful of early B-type stars (e.g. Donati et al. 2001; Neiner et al. 2003; Alecian et al. 2008; Petit et al. 2008; Silvester et al. 2009). Indications of magnetic fields have also been reported in a number of other OB stars (e.g. by Hubrig et al. 2008; 2009). Such objects represent vital targets for the study of stellar magnetism. Their strong, radiatively-driven winds couple to magnetic fields, generating complex and dynamic magnetospheric structures (e.g. ud-Doula \\& Owocki 2002), which modify mass loss, and may enhance the shedding of angular momentum via magnetic braking (e.g. ud-Doula et al. 2008). The presence of even a relatively weak magnetic field can profoundly influence the evolution of massive stars and their feedback effects, such as mechanical energy deposition in the ISM and supernova explosions (e.g. Ekstrom et al. 2008). For these reasons, the Magnetism in Massive Stars (MiMeS) Large Program (Wade et al. 2009; Grunhut et al. 2009) is exploiting the unique spectropolarimetric characteristics of ESPaDOnS at the Canada-France-Hawaii Telescope (CFHT) and NARVAL at the T{\\' e}lescope Bernard Lyot (TBL) to obtain critical missing information about the poorly-studied magnetic properties of these important stars.  One of the MiMeS targets is HD~57682, an O9 sub-giant star. HD~57682 is reported to be a runaway star (Comeron et al. 1998) with no known companions (e.g. de Wit et al. 2005; Turner et al. 2008) and shows no photometric variability (e.g. Balona 1992). However, HD~57682 exhibits variability of UV lines characteristic of magnetic OB stars (e.g. Schnerr et al. 2008), and this is the primary reason it was included in the MiMeS target list. ",
        "conclusions": "\\label{disc_sec} It is interesting to note that the mass-loss rate derived from the H$\\alpha$ profiles is inconsistent with mass-loss derived from the C~{\\sc iv} UV lines. Test models indicate that the intensity of the H$\\alpha$ emission observed at different dates requires mass-loss rates above $\\sim$$10^{-7}$~M$_\\odot$ yr$^{-1}$. With such high values, the predicted UV spectra would present strong P-Cygni profiles (e.g. Marcolino et al. 2009), which are not seen in any of the IUE spectra available for HD~57682. In fact, HD~57682 presents a UV spectrum typical of Galactic weak wind stars, with no significant P-Cygni emissions (Martins et al. 2005b; Marcolino et al. 2009). This leads us to believe that the H$\\alpha$ profiles are not formed in the wind, but are possibly a result of confined circumstellar material. Walborn (1980) already reported this line to be unusual and of non-nebular origin.  A preliminary Bayesian analysis suggests that the magnetic field of HD~57682 is approximately reproduced by a dipole field with a characteristic surface field strength of $1680~^{+134}_{-356}$~G. We computed the magnetic wind confinement parameter (ud-Doula \\& Owocki 2002) $1.4\\times10^4$ for HD~57682, using $B_p\\sim1600$~G, and the physical parameters listed in Table~\\ref{params} ($\\eta_*$ ranges from $4\\times10^3$ - $2\\times10^4$ using the 99.7\\% credible regions). Although strongly confined, this places HD~57682 at an interesting intermediate regime amongst OB stars, between the super-strongly-confined wind of $\\sigma$~Ori~E ($\\eta_*\\sim10^7$) and the much more weakly confined wind of $\\theta^1$~Ori~C ($\\eta_*\\sim20$). Magnetic confinement of the wind of HD~57682 would naturally explain the H$\\alpha$ variability, the UV variability, and the line depth variability. Moreover, shocks produced in a magnetically channelled wind could boost the X-ray emission, providing an explanation of the high X-ray luminosity without requiring a high mass-loss rate (e.g. Cohen et al. 2008). However, it is puzzling that HD~57682 presented a soft X-ray spectrum during the ROSAT observations, contrary to what would be expected from a fast wind and the strong magnetic confinement.  This discovery brings the number of firmly-established magnetic O-type stars to four. With its strong wind confinement, and its stellar and wind properties determined, HD 57682 provides a testbed for future MHD simulations. However, more data are needed to better determine the surface dipole field strength, magnetic obliquity, and inclination."
    },
    "0910/0910.0022_arXiv.txt": {
        "abstract": "With the coming generation of instruments and telescopes capable of spectroscopy of high redshift galaxies, the spectral synthesis technique in the rest-frame UV and Far-UV range will become one of a few number of tools remaining to study their young stellar populations in detail. The rest-frame UV lines and continuum of high redshift galaxies, observed with visible and infrared telescopes on Earth, can be used for accurate line profile fitting such as P\\,{\\sc{v}}$\\lambda\\lambda$1118,1128, C\\,{\\sc{iii}}$\\lambda$1176, and C\\,{\\sc{iv}}$\\lambda$1550. These lines are very precise diagnostic tools to estimate ages, metallicities, and masses of stellar populations.  Here we discuss the potential for spectral synthesis of rest-frame UV spectra obtained at the Keck telescope. As an example, we study the 8 o'clock arc, a lensed galaxy at z=2.7322. We show that the poor spectral type coverage of the actual UV empirical spectral libraries limits the age and metallicity diagnostic. In order to improve our knowledge of high redshift galaxies using spectral synthesis, UV stellar libraries need to be extended to obtain accurate age, metallicity, and mass estimates likely to be occuring in young stellar populations observed in the early universe. ",
        "introduction": "The study of young stellar populations has benefited greatly from the ultraviolet (UV) evolutionary spectral synthesis technique, well-developed in the past two decades. In the 1000-2000\\,\\AA\\ range, many lines display  P\\,Cygni profiles typical of strong and dense stellar winds. The work by \\cite{rob93} and \\cite{lei95} have used IUE and HST/STIS UV spectra to show how the UV line profiles such as Si\\,{\\sc{iv}}$\\lambda$1400, C\\,{\\sc{iv}}$\\lambda$1550, and He\\,{\\sc{ii}}$\\lambda$1640 were sensitive to the spectral type, luminosity class, and metallicity of early-type stars. The authors made used the line profiles observed in IUE and HST/STIS spectra to build an empirical stellar library for stellar population evolutionary models. This tool opened the door to the precise determination of the main physical properties of young stellar populations such as age, stellar mass, metallicity, and mass function (\\cite[Leitherer \\etal\\ 1999]{lei99}). Later on, the UV range below Ly$\\alpha$ became accessible with FUSE and more UV features were added to the collection of line profiles useful for evolutionary spectral synthesis (\\cite[Pellerin \\etal\\ 2002]{pel02}, \\cite[Robert \\etal\\ 2003]{rob03}). Good diagnostic lines found in this regime are C\\,{\\sc{iii}}$\\lambda$1176, and P\\,{\\sc{v}}$\\lambda\\lambda$1118, 1128.  So far the use of UV line profiles has proven to be quite successful to study young stellar populations in nearby galaxies (e.g. \\cite[Pellerin 2006]{pel06}, \\cite[Pellerin \\& Robert 2007]{pel07}), and the future of UV spectral synthesis is expected to be quite bright as well. Indeed, the rest-frame UV spectra of high-redshift galaxies are now accessible with current optical facilities and many more of them will be observable with the coming generation of 30 meter class telescopes. ",
        "conclusions": ""
    },
    "0910/0910.2027_arXiv.txt": {
        "abstract": "{ Isotropic diffusion models for Galactic cosmic ray transport put tight constraints on the maximum convection velocity in the halo. For a half halo height of 4 kpc the maximum convection speed is limited to 40 km/s in the halo, since otherwise the constraints from local secondary to primary ratios and radioactive isotopes cannot be met. The ROSAT Galactic wind observations of wind speeds up to 760 km/s therefore constitute a problem for diffusion models.\\\\ It is shown that such wind speeds are possible, if the diffusion coefficient in the halo is different from the diffusion coefficient in the disk. The radial dependence of the wind velocity was taken to be proportional to the source strength, as expected from winds which are sustained by cosmic ray pressure. In this case the cosmic ray density and with it the diffuse $\\gamma$-ray production from nuclear interactions are suppressed near the sources. This solves in a natural way the problem of the soft gradient in the radial dependence of the $\\gamma$-ray flux. Furthermore, the large bulge over disk ratio in positron annihilation as observed by INTEGRAL, can be explained by positron escape from the disk in such a model.} ",
        "introduction": "High energy cosmic rays (CRs) are thought to be accelerated in the shock waves of supernovae remnants (SNRs), as can be found in standard text books \\citep{Ginzburg:1990sk,Longair,Schlickeiser:2002pg}, but critical reviews exist \\citep[see e.g. ][]{kirk:2001}. Although CRs are accelerated to almost the speed of light they are confined to our Galaxy for about $10^7$ yrs, as is known from the presence of radioactive nuclei with decay times of a few times $10^6$ yrs, see e.g. \\citep{Cesarsky:1980}. Such ''cosmic clocks'' require averaged drift speeds of CRs of only  a fraction of the speed of light. This is possible if CRs continuously scatter on magnetic turbulences \\citep{Ginzburg:1990sk} and as a consequence perform a random walk which can be modeled by diffusion. Possible scattering centers are the turbulent magnetic fields generated by the dilute plasma of CRs themselves. In the simplest case diffusion is assumed to be isotropic with the same diffusion coefficient in the halo and the disk \\citep{Strong:2007nh}. The parameters of the diffusion equation can be obtained by taking a source distribution proportional to the SNR distribution and measuring the local density and energy spectra of the CRs, which depend on the transport parameters, gas densities and magnetic fields between the source and the solar system.  An isotropic diffusion model has been implemented in the publicly available GALPROP code\\footnote{The GALPROP code is publicly available under http://galprop.stanford.edu}  \\citep{Strong:1998pw, Moskalenko:1998id}. Details about the underlying propagation and the parameter determination can be found in a review by \\citet{Strong:2007nh}.  In addition to diffusion convective transport modes may play a role for CRs \\citep{Jokipii:1976jk}. Supernovae (SNs) eject hot gas into the interstellar medium (ISM) which can expand into the halo, presumably driven by the CR pressure from supernovae remnants (SNRs)\\citep{Breitschwerdt:2008na}. This  leads to a reduction of the CR density and CR interaction rate in the Galactic disk, most efficiently at radii with high source density. This mechanism has been proposed as an explanation for the relatively small  production of diffuse $\\gamma$-rays near the sources in comparison to the $\\gamma$-ray production at larger radii \\citep{Breitschwerdt:2002vs}, known as the ''soft $\\gamma$-ray gradient (SGRG) problem''. \\\\ Unfortunately, the maximum allowed convection speed in isotropic propagation models is restricted to a few tens of km/s \\citep{Strong:2007nh}, because otherwise the constraints from local secondary to primary ratios and radioactive isotopes cannot be met. Wind velocities as small as this will not solve the soft-$\\gamma$-ray-gradient problem. Within the isotropic diffusion models alternative solutions have been investigated; e.g. putting more sources in the outer Galaxy \\citep{Strong:1998pw} or putting more gas there by assuming that the tracer for molecular hydrogen, the radiation from CO molecules, has a different proportionality factor in the outer Galaxy \\citep{Strong:2004td}, as motivated by the change in metallicity. A recent analysis of  the ROSAT data on x-rays implies wind speeds of up to 760 km/s in the halo \\citep{Everett:2007dw,Breitschwerdt:2008na}. Speeds like this are sufficient to solve the SGRG problem, but they would require to loosen the constraint of isotropic diffusion.\\\\  An additional problem for isotropic diffusion models, is the large bulge-over-disk (B/D) ratio of the 511 MeV positron annihilation line, as observed by the INTEGRAL satellite \\citep{Knodlseder:2005yq,Weidenspointner:2007rs}. These low energy positrons are thought to originate predominantly from the decay of radioactive nuclei, produced in SNIa explosions in the bulge and in the disk. Predictions of the B/D ratio from the expected number of SN explosions are well below one, because of the high rate of SNIa explosions in the disk \\citep{Prantzos:2005pz}. However, INTEGRAL found a B/D ratio of a few. In a diffusion model without convection MeV positrons hardly propagate, because diffusion is proportional to the velocity and energy of the particle. In this case positrons annihilate close to their sources leading to a small B/D ratio. Convective transport in Galactic winds is independent of energy, so low energy positrons in the disk can be convected to the halo, where there are hardly electrons at rest to annihilate with. In this case the escape fraction of  positrons from the disk can be sufficient to explain the observed large B/D ratio.  In this paper an anisotropic transport model, which allows for ROSAT compatible convection and still meets all the constraints from primary and secondary cosmic rays, cosmic clocks, $\\gamma$-rays and the INTEGRAL B/D ratio, is introduced. The model features different diffusion coefficients in the halo and in the disk. By increasing the diffusion coefficient towards the halo boundary a smooth transition to free escape is obtained, so the model becomes insensitive to the precise position of the boundary. This is in strong contrast to isotropic propagation models, which require a precise size of the halo for a given diffusion constant.  The paper is organized as follows: in section \\ref{s_iso} we summarize the main features of the current isotropic diffusion models for CR transport. Section \\ref{s_limits} discusses the limits of the isotropic transport models in the light of observational constraints which refine our picture of CR transport. In section \\ref{s_aniso} we introduce the basic ideas of an anisotropic propagation model (aPM). Section \\ref{s_perf} discusses the parameter determination of the aPM. The performance of the aPM with respect to charged CRs and diffuse Galactic $\\gamma$-rays is discussed in sections \\ref{s_perf_CR} and \\ref{s_perf_gamma}, respectively. Section \\ref{s_conc} summarizes the results. ",
        "conclusions": "\\label{s_conc} A new model for Galactic cosmic ray transport is presented, which allows for significant convective transport as expected from the Galactic winds deduced from the x-ray data from the ROSAT satellite. The model has been realized by modifying the publicly available GALPROP code, which up to now allowed only isotropic transport. Isotropic transport models, which feature globally constant transport parameters, can only accommodate negligible convection speeds, since with convection the CRs are driven away from the disc, thus  not returning often enough to the disk to produce secondary particles.  A possible way to allow for large convection is to allow the diffusion in the halo to be different from the diffusion in the disk, i.e. to give up isotropic propagation with globally constant transport parameters. The GALPROP code was modified in the following way: \\begin{itemize} \\item the Galactic winds were assumed to be proportional to the CR source distribution, which was taken to be the SNR distribution \\item the mean free path of CRs - and therefore the diffusion coefficient - in the halo was assumed to increase linear with the distance from the disk. Such a choice is motivated  if the scattering centers are produced by the CRs themselves (via the irregular magnetic field components in the plasma) and the density of CRs drops approximately linearly in the halo. \\end{itemize} Fixing the magnitude of the convection speed to the wind speeds suggested by the ROSAT data, the increase in diffusion coefficient in the halo can be fitted from the amount of secondary production (from B/C ratio) and the residence time of CRs (from the cosmic clocks, in this case the $^{10}Be/^9Be$ ratio). It is shown that such a  model is consistent with all available CR data, including not only the ROSAT data on convective winds, but also the large bulge/disk ratio of the positron annihilation line as observed  by the INTEGRAL satellite. In isotropic propagation models one has to introduce a new source of MeV positrons in the bulge, like e.g. dark matter annihilation or tunneling the positrons from the disk to the bulge via regular magnetic fields in the halo. In the anisotropic model presented here the large B/D ratio is naturally explained by the energy independent convective transport of low energy positrons from the disk to the halo, where there are no electrons to annihilate with. Convective transport is absent in the bulge, because the gravitational potential is too strong there to launch Galactic winds.  An additional interesting feature of the present model is the smooth transition to free escape of CRs, because of the increase in mean free path with increasing distances from the disk. Therefore the boundary condition can be moved to infinity in contrast to isotropic propagation models, where the boundary condition is fine-tuned to get the correct residence time of CRs inside the Galaxy."
    },
    "0910/0910.5741_arXiv.txt": {
        "abstract": "Gamma-ray emission from pulsars has long been modeled using a vacuum dipole field. This approximation ignores changes in the field structure caused by the magnetospheric plasma and strong plasma currents. We present the first results of gamma-ray pulsar light curve modeling using the more realistic field taken from three-dimensional force-free magnetospheric simulations. Having the geometry of the field, we apply several prescriptions for the location of the emission zone, comparing the light curves to observations. We find that when the emission region is chosen according to the conventional slot-gap (or two-pole caustic) prescription, the model fails to produce double-peak pulse profiles, mainly because the size of the polar cap in force-free magnetosphere is larger than the vacuum field polar cap. This suppresses caustic formation in the inner magnetosphere. The conventional outer-gap model is capable of producing only one peak under general conditions, because a large fraction of open field lines does not cross the null charge surface. We propose a novel ``separatrix layer\" model, where the high-energy emission originates from a thin layer on the open field lines just inside of the separatrix that bounds the open flux tube. The emission from this layer generates two strong caustics on the sky map due to the effect we term ``Sky Map Stagnation\" (SMS). It is related to the fact that the force-free field asymptotically approaches the field of a rotating split monopole, and the photons emitted on such field lines in the outer magnetosphere arrive to the observer in phase. The double-peak light curve is a natural consequence of SMS. We show that most features of the currently available gamma-ray pulsar light curves can be reasonably well reproduced and explained with the separatrix layer model using the force-free field. Association of the emission region with the current sheet will guide more detailed future studies of the magnetospheric acceleration physics. ",
        "introduction": "Gamma-ray pulsars are of great scientific interest in high-energy astrophysics. The high energy emission from these objects provides valuable information about pulsar magnetospheric structure, particle acceleration mechanisms, and plasma physics in strong magnetic fields. Observations by the Energetic Gamma Ray Experiment Telescope ({\\it EGRET}) onboard Compton Gamma Ray Observatory ({\\it CGRO}) confirmed 7 gamma-ray pulsars \\citep{Thompson04}. The light curves of these pulsars are characterized by widely separated double-peak features, with the first peak typically lagging the radio pulse by a small fraction of rotation period. Recently, {\\it AGILE} mission has detected several new gamma-ray pulsars and refined the detection of some previously known pulsars \\citep{Halpern_etal08,AGILE09a,AGILE09b}. More impressively, one year after launch, the Large Area Telescope ({\\it LAT}) onboard {Fermi Gamma-ray Space Telescope} ({\\it Fermi}) has discovered more than 40 new gamma-ray pulsars \\citep{FermiCTA1,Fermi09a,Fermi09b, FermiRelease09a,FermiRelease09b,FermiPulsars09}. The sensitivity and timing capability of {\\it Fermi-LAT} have also provided the most precise light curves and phase resolved pulsar spectra \\citep{FermiVela, Fermi09c,Fermi09d}. These new discoveries have greatly expanded the sample of gamma-ray pulsars and underscored the importance of understanding the origin of double-peak profiles. At the same time, the diversity of light curves from latest gamma-ray pulsar sample calls for improvements in theoretical modeling.  The widely separated double peak feature in the gamma-ray light curves suggests emission coming from the outer magnetosphere. Gamma rays are believed to be produced by energetic particles accelerated in the ``gap\" regions in the magnetosphere. The particles emit via curvature, synchrotron and inverse Compton (IC) radiation. Various theoretical models differ in the assumed location of the emission zones (i.e., the gaps). Conventional models for pulsar gamma-ray emission include the polar cap model \\citep{Harding_etal78,DH82,DH96}, the slot gap model (SG, or two-pole caustic model, TPC for short; \\citealp{AS79,Arons83,MH03,MH04a,DyksRudak03,DHR04}), the outer gap model(OG, \\citealp{CHR86a,CHR86b,RY95,Yadi97,CRZ00}), as well as the inner annular gap model (IAG, \\citealp{Qiao_etal04,Qiao_etal07}).  Most calculations from these models approximate the pulsar magnetosphere by a rotating magnetic dipole field. In reality, the magnetosphere is known to be filled with plasma \\citep{GJ69}. This plasma is essentially force-free (FF), satisfying $\\rho\\mb{E+j\\times B}/c=0$, where $\\rho$ and ${\\mb j}$ are the charge and current densities. The structure of the FF magnetosphere differs substantially from the rotating dipole field due to the poloidal current and the returning current sheet \\citep{CKF99,Gruzinov05,Komissarov06,McKinney06,Timokhin06}. Recently, FF magnetospheric field structure in 3D has become available from the time-dependent FF simulations \\citep{Spitkovsky06,KC09}. The FF magnetosphere satisfies ${\\mb E\\cdot \\mb B}=0$ everywhere, hence no particle acceleration or emission can formally occur. Real pulsars must lie somewhere between the vacuum, which has ${\\mb E\\cdot \\mb B}\\ne0$ everywhere, and FF cases. The fact that the power of gamma-ray emission is a small, though non-negligible ($\\lesssim 20\\%$), fraction of the total spin down power suggests that FF should be a reasonable approximation to the overall field structure with the exception of relatively small regions where acceleration takes place.  In the companion paper (\\citealp{BS09a}, hereafter BS10), we demonstrated that the modeling of gamma-ray pulsar light curves using the vacuum dipole field is subject to large uncertainties. This is because the peaks in the vacuum light curve are caused by the fortuitous overlap of radiation from different regions of the magnetosphere (caustics). This overlap is sensitive to the geometry of the magnetic field due to two effects. First, the shape of the field controls how the aberration of light and light travel delay add together to cause the formation of caustics. We showed that a seemingly small change of treating the retarded vacuum field in the instantaneously corotating frame, rather than in the laboratory frame where it should be defined, can lead to significant reduction in the strength of caustics in the two-pole caustic/slot gap model and to the moderate dilution of the outer gap peaks. Second, the behavior of the field near the light cylinder determines the shape of the polar cap on the star, and, hence, indirectly controls the shape of the emission zone, even if it is located in the inner magnetosphere. We considered the sensitivity of the caustic formation to the modification of shape of the vacuum polar cap. We compared the polar cap obtained by tracing the vacuum field lines to a simple circular polar cap. This change significantly affected the caustics of the two pole caustic model. The outer gap model was less sensitive to this change, but, since the emission zone for the outer gap is in the outer magnetosphere, the field shape there is likely to be strongly modified by the inclusion of plasma effects. Given these uncertainties in the vacuum field modeling, it is highly desirable to revisit the existing models using the FF field configuration and compare the results with the latest observations.  In this paper, we use the force-free field from 3D time-dependent simulations by \\citet{Spitkovsky06} (hereafter, S06), and present the calculation of pulsar gamma-ray light curves using various theoretical models of emission. In \\S2, we provide the general analysis of FF field structure, addressing the location of the current sheet and its association with magnetic field lines. In \\S3, we describe our numerical method and theoretical models for calculating the gamma ray light curves. \\S4-\\S6 are dedicated to comparison between different models. We show in \\S4 that the conventional two-pole caustic model fails to produce double-peak light curves using the FF field. We propose a novel ``separatrix layer\" (SL) model in \\S5, in which the emission zone is fixed in a layer just inside the separatrix associated with the strong current sheet. This layer can be accurately traced by open field lines originating in an annulus just inside the polar cap rim\\footnote{In the original version of this manuscript, we termed this model as ``annular gap\" model. The name ``separatrix layer\" more accurately describes the location of the emission zone and avoids the confusion with the ``inner annular gap\" model \\citep{Qiao_etal04}. Moreover, although the present paper is based only on geometry, the association of the emission layer with the magnetospheric current sheet (see \\S\\ref{ssec:origin}) implies that the acceleration may not necessarily involve a ``gap\" at all. }. This model differs from the conventional TPC model in that the emission zone is not concentrated on the last open field lines, but rather interior to them, and it spans a larger range of heights up to and beyond the light cylinder. It differs from the conventional OG model in that the emission zone consists of all field lines in a given flux tube, and not just those that cross the null charge surface. The double-peak light curve is a generic property of the FF field structure, where caustics of emission form in the outer magnetosphere when the field lines approach the split-monopolar geometry. We call this ``the stagnation effect\" on the sky map. In \\S6, we show that the conventional OG model applied to FF field has difficulty in producing double-peak light curves unless a carefully chosen inclination angle and viewing geometry is used. We discuss various implications of our results to the pulsar acceleration mechanisms and compare our results with recent observations in \\S7. The main results are summarized in \\S8. ",
        "conclusions": "We have shown that using the FF field, conventional TPC and OG model no longer work as well as in previous calculations using the vacuum dipole field. Alternatively, we find that the separatrix layer model, where emission comes from open field lines that lie just in the vicinity of the separatrix (including the strong current layer inside the LC and the current sheet outside the LC), is very promising in reproducing the prevailing double peak features of gamma-ray pulsar light curves. This model works due to the sky map stagnation effect, that is unique to FF field and occurs as the field becomes increasingly split-monopolar with radius. Our calculations of the SL model in \\S\\ref{sec:AG} work mainly on the geometric grounds with little physical input. In this section we discuss the implications of these results to pulsar gamma-ray radiation mechanism, and compare our results with observational data.  \\subsection[]{Applicability of the FF field}\\label{ssec:FFapply}  The FF field provides a more realistic model for the pulsar magnetosphere than vacuum field because it takes into account the field distortions due to self-consistent currents caused by pulsar spin down. The basic assumption of the force-free approximation is that pair creation is so efficient that it fills the magnetosphere with abundant pair plasma to make it a perfect conductor everywhere. Real pulsar magnetospheres are more complicated due to dissipation and reconnection in the current sheet, formation of gaps, and the presence of differential rotation \\citep{Timokhin07a,Timokhin07b}.  One way of testing the robustness of the result obtained in this paper is to consider fields from  strong-field electrodynamics (SFE) introduced by \\citet{Gruzinov07a,Gruzinov08a,Gruzinov08b}, which generalizes the FF equations to include dissipation. In SFE, the FF condition ${\\mb E}\\cdot{\\mb B}=0$ is relaxed in space-like regions, where the Ohmic dissipation is present, while time-like regions are still non-dissipative. The dissipation and reconnection in the FF current sheet may be better modeled with SFE. In addition, vacuum gaps in pulsar magnetosphere may also be considered as dissipative regions, where ${\\mb E}\\cdot{\\mb B}\\neq0$. We have implemented SFE formulation into our FF code by replacing the FF current by the SFE current [see equation (6) of \\citet{Gruzinov07b}]. The conductivity scalar $\\sigma=\\sigma(E_0,B_0)$ is taken to be constant. The natural scale of $\\sigma$ is $\\Omega/c^2$, which we denote by $\\sigma=1$. We have run our simulation with different values of $\\sigma$ around 1. We find that the overall field configuration is more smooth than the FF field configuration. Particularly, the ``knot\" in the sky map from LOFLs which penetrate the current sheet becomes less prominent. The strong current layer inside the LC is not changed qualitatively, but the current sheet outside the LC is smoothed appreciably. More energy is dissipated in the magnetosphere, and the total Poynting flux decreases by  about $10\\%$ over about $1R_{LC}$.  Despite the fact that SFE produces smoother magnetospheric structure, we find that the resulting sky maps and light curves do not differ qualitatively from the results with the FF field presented in this paper. Therefore, we believe that our result is robust and may also be applicable to dissipative magnetospheres with moderate levels of dissipation.   \\subsection[]{The Origin of the Separatrix Layer Emission}\\label{ssec:origin}  We have shown that if pulsar high-energy emission originates from open field line regions in the vicinity of the separatrix, extending from the NS surface to beyond the LC, two bright caustics will form due to sky map stagnation. Currently, there is no theory directly pointing to the existence of the SL emission; however, based on the geometric location of the SL, there are clues about its origin.  1. Modified SG/TPC model. We note that SMS is already significant at $r_{\\rm ov}=0.95$, which is very close to LOFLs. Also, in our calculation, the emission zone is assumed to have a Gaussian profile with width $\\Delta r_{\\rm ov}=0.025$ centered on $r_{\\rm ov}$. Therefore, the edge of the emission zone can be considered on the LOFLs. This emission zone geometry is consistent with the SG model. Therefore, if the thickness of the pair formation front is as large as $\\Delta r_{\\rm ov}=0.025-0.05$, a modified SG model may be able to produce the desired two peak light curves. Such SG must extend at least to the LC to produce sharp peaks. Recent SG/TPC models consider emission from smaller $r_{\\rm ov}$ flux tubes (e.g., \\citealp{Venter_etal09}). Given that in FF field only the field lines with $r_{\\rm ov}<0.95$  are truly open and do not close through the current sheet, it is likely that at least geometrically the SL model can be considered as an extension of TPC and invokes the same field lines.  The physical reason for the acceleration on these lines is unlikely to be a slot gap, however, and is more  related to reconnection in the current sheet, as discussed below.  2. Modified OG model. Another clue comes from the distribution of space-like regions in the magnetosphere. Regions of space-like 4-current, which are likely to develop instabilities and could have accelerating fields \\citep{Gruzinov07b}, are preferentially located beyond the NCS. Similar locations are typically invoked in OG models, although for reasons that rely on complete charge separation in the magnetosphere. Even if OG-type model can be justified based on the sign of 4-current, it would not be sufficient to produce two peaks -- the emission has to come from more field lines than those crossing the NCS. Interestingly, looking at Figure \\ref{fig:AGSM60_90}d, the regions showing the pile-up due to SMS have values of $\\varrho$ near zero, or negative. If small positive values of $\\varrho$ can also cause acceleration, this could explain the emission from the whole flux tube.  3. Association with the strong current. The presence of a strong current layer within the LC and a current sheet beyond LC are general features of FF field. Strong dissipation and reconnection is expected in such regions \\citep{Gruzinov07a,Lyubarsky08}, and the energy release can reach a few percent of the spin down power, probably radiated in X-rays and gamma rays \\citep{Lyubarsky96}. Using Bogovalov's (1999) current sheet solution for the inclined split-monopole, \\citet{Kirketal02} were able to reproduce the typical double-peak light curves, indicating that the field configuration around the current sheet beyond the LC is also geometrically favored. Based on the same model, \\citet{PetriKirk05} and \\citet{Petri09} further reasonably well reproduced the optical polarization of the Crab pulsar and phase-resolved spectrum for the Geminga pulsar. Our result that the SL emission zone resides in the vicinity of the strong current layer and/or current sheet (see Figure \\ref{fig:cusheet}) implies a connection with the strong current. At higher inclinations, the return current becomes increasingly distributed over the open flux, and it is possible that some of the acceleration takes place on the open field lines. Also, the caustics of emission form in the outer magnetosphere, and the acceleration on these field lines can be affected by the proximity of the Y-point. Reconnection near the Y-point (or ``Y-ring\" in 3D) can load the open field lines with plasma, and cause field-aligned accelerating fields, akin to auroral double-layers (\\citealp{Arons08}). Emission from the equatorial current sheet beyond the LC can also be important, as the field lines we invoke for SL  trace close to the current sheet in the wind zone. Modeling emission from the current sheet itself is more complicated in FF simulation, since ideal MHD doesn't apply in the current sheet, and the magnetic field direction abruptly changes in the unresolved current sheet, which leads to tracing errors. This will be studied more in future work. The facts that the bulk of the SL emission is formed at and beyond the LC, and that emission from two poles overlaps on the sky map and forms two peaks only close to the edge of the open flux tube (for $r_{\\rm ov} > 0.9$, Figure \\ref{fig:rov_cmp}), suggests that the current sheet beyond the LC is responsible for the large part of gamma-ray emission.  Our discussion on the origin of the SL model is highly speculative. A detailed theory of the electrodynamics of current sheets has to be constructed  to test these scenarios and uncover the origin of the SL emission.  \\subsection[]{Comparison with Observations}\\label{ssec:observation}  One year after launch, {\\it Fermi} LAT has greatly expanded the number of detected gamma-ray pulsars either by identifying previously unresolved/unidentified sources (e.g., known radio pulsars, millisecond pulsars (MSP), EGRET sources, supernova remnants (SNR), pulsar wind nebulae (PWN), \\citealp{FermiCTA1,Fermi09a,Fermi09b,FermiRelease09b}), or by the blind search of LAT sources \\citep{FermiRelease09a}. Together with previously known gamma-ray pulsar data, a sample of more than 40 gamma-ray pulsars is now available \\citep{FermiPulsars09}, allowing more systematic comparison between observations and theoretical models.  Under the framework of the SL model with FF field, the double-peak light curves can be produced over a large range of geometric parameters, as shown in \\S\\ref{sec:AG}. As a geometric model, it provides robust prediction on the temporal location of the two peaks regardless of the input physics. The phase lag of the first peak relative to the radio peak (which likely coincides with the location of the polar cap on the sky map) ranges from about $\\delta=0.14$ to larger than $\\delta=0.3$, and the separation between the two peaks ranges from zero (i.e., one peak) to $\\delta=0.5$. Smaller phase lag typically corresponds to larger separation, and vice versa, as seen from the caustic structure from Figure \\ref{fig:AGLC_10}. Also, larger inclination angles tend to result in larger peak separation. Statistically, our model predicts that the majority of light curves should have widely-separated double peaks, pulsars with smaller peak separation are less common, pulsars with only one peak are possible but rare. Another consequence of the SL model is that smaller peak separation usually indicates stronger bridge emission between the peaks. The best example of this can be drawn from the light curve with $\\alpha=45^\\circ$, $\\xi_{\\rm obs}=50^\\circ$ in Fig. \\ref{fig:AGLC_10}.  Based on the current data sample of gamma-ray pulsars, the majority of them show widely separated double-peak pulse profiles, in good agreement with the SL model. A small fraction shows double peaks with smaller separation [e.g., B1716-44 \\citep{AGILE09a}, J0007+7303, J1459-60, J1741-0254 \\citep{FermiRelease09a}], but with relatively strong bridge emission, also in qualitative agreement with our result\\footnote{Note that in Figure 2 of \\cite{FermiRelease09a}, the first peak is placed at phase of 0.3 because most of the new gamma-ray pulsars are radio quiet.}. A few objects show only one peak [e.g., J0357+32 \\citep{FermiRelease09a}, PSR 0437-4715, PSR 0613-0200 \\citep{FermiRelease09b}, J2229+6114 \\citep{AGILE09b}]. Such pulse profiles may be obtained when the observer's line of sight tangentially cuts through the caustics on the sky map (examples include $\\alpha=30^\\circ, \\xi_{\\rm obs}=50^\\circ$ in Figure \\ref{fig:AGLC_10}). The small number of such pulsars agrees with the small chance of such tangential incidence. Noticeably, the fraction of single peak MSPs appears to be higher than for normal gamma-ray pulsars \\citep{FermiRelease09b}. This may reflect the fact that MSPs tend to have smaller inclination angle which increases the chance of tangential incidence. Given the relatively old age of the MSP population, their smaller inclination angle might be caused by alignment torques \\citep{DavisGoldstein70,Goldreich70,Melatos2000}.  The distribution of phase lags from the SL model is not directly comparable with the latest data since most new gamma-ray pulsars are radio-quiet. For available radio-loud pulsars, we find some agreement, but not always. For example, the phase separation between peaks in B1706-44 appears to be too small to fit our model, as both peaks come in the first half of the rotation. There may be an additional phase lag from the radio peak due to high altitude radio emission from energetic pulsars \\citep{WeltevredeJohnston08}. Therefore, it is uncertain whether observed phase lag can be used as a reliable diagnostics."
    },
    "0910/0910.0678_arXiv.txt": {
        "abstract": "We report theoretical and experimental results of on-going feasibility studies to detect cosmic neutrinos acoustically in Lake Baikal. In order to examine ambient noise conditions and to develop respective pulse detection techniques a prototype device was created. The device is operating at a depth of 150 m at the site of the Baikal Neutrino Telescope and is capable to detect and classify acoustic signals with different shapes, as well as signals from neutrino-induced showers. ",
        "introduction": "During the last years, neutrino astronomy is undergoing rapid development and becomes a new ``window'' to the Universe.  The large scale neutrino telescopes currently under operation (NT200+\\,\\cite{NT200, NT+} in Lake Baikal, AMANDA/IceCube \\cite{IceCube} at the South Pole and ANTARES \\cite{ANTARES} in the Mediterranean) detect the Cherenkov light emitted in water or ice by relativistic charged particles produced via neutrino interactions with matter.  Back in 1957, G.A. Askaryan has shown that a high-energy particle cascade in water, besides the Cherenkov radiation, should also produce an acoustic signal \\cite{Askarian-1957}. The potential of the acoustic detection of particle cascades is based on the fact that the absorption length for acoustic waves with a frequency about $30$ kHz in sea water is at least an order of magnitude larger than that of Cherenkov radiation, in the fresh Baikal water this ratio is even close to 100 \\cite{Clay}. The second fact that is favorable for the detection of acoustic signals from cascade showers at distances of hundreds of meters, or even at such long distances as several kilometers, is that the amplitude of pulses produced by showers in the near-field zone decreases only as the square root of distance, while in the far-field zone it decreases as the reciprocal of distance from the shower \\cite{Askarian-1979, Learned}. Therefore, in principle, a deep-water acoustic detector of high-energy neutrinos can have a much smaller number of measuring channels compared to a Cherenkov detector with the same effective volume \\cite{Askarian-1977}.  However, the technology of acoustic detection in high-energy physics is much less developed than optical methods. Since several years, however, an increasing number of feasibility studies on acoustic particle detection are performed [10--16].  Actually, the possibility of acoustic detection of high-energy neutrinos and the energy detection threshold are determined by the possibility of separating the signals produced by cascade showers from noise produced by other sources.\\\\ \\begin{figure*}[ht] \\centering \\includegraphics [width=0.88\\textwidth]{icrc0927_fig01.eps} \\caption{Schematic view of the underwater $4$-channel digital device for detection of acoustic signals from high energy neutrinos.} \\label{device} \\end{figure*} ",
        "conclusions": ""
    },
    "0910/0910.1379.txt": {
        "abstract": "Based on the physical model of a supermassive black hole (SMBH) growth via gas accretion in a circumnuclear disk (CND) proposed by Kawakatu \\& Wada (2008), we describe the formation of high-$z$ ($z > 6$) quasars (QSOs) whose BH masses are $ M_{\\rm BH}> 10^{9}M_{\\odot}$. We derive the necessary conditions to form QSOs at $z > 6$ by only gas accretion: %in terms of the total accreted gas mass from host galaxies, %$M_{\\rm sup}$ and star formation efficiency in the CND. %As a result, we find that the following conditions are required; (i) A large mass supply with $M_{\\rm sup} > 10^{10}M_{\\odot}$ from host galaxies to CNDs, because the final BH mass is only $1-10\\%$ of the total supplied mass from QSO hosts. (ii) High star formation efficiency for a rapid BH growth which is comparable to high-$z$ starburst galaxies such as submillimeter galaxies (SMGs). We also find that if the BH growth is limited by the Eddington accretion, the final BH mass is greatly suppressed when the period of mass-supply from hosts, $t_{\\rm sup}$ is shorter than the Eddington timescale. Thus, the super-Eddington growth is required for the QSO formation as far as $t_{\\rm sup}$, which is determined by the efficiency of angular momentum transfer, is shorter than $\\sim 10^{8}\\,{\\rm yr}$. % The evolution of the QSO luminosity depends on the redshift $z_{\\rm i}$ at which accretion onto a seed BH is initiated. In other words, the brighter QSOs at $z >6 $ favor the late growth of SMBHs (i.e., $z_{\\rm i}\\approx 10$) rather than early growth (i.e., $z_{\\rm i} \\approx 30$). For $z_{\\rm i}\\approx 10$, $t_{\\rm sup}\\simeq 10^{8}\\,{\\rm yr}$ is shorter than that of the star formation in the CND. Thus, the gas in the CND can accrete onto a BH more efficiently, compared with the case for $z_{\\rm i}\\approx 30$ (or $t_{\\rm sup}\\approx 10^{9}\\,{\\rm yr}$). Moreover, we predict the observable properties and the evolution of QSOs at $z >6$. In a QSO phase, there should exist a stellar rich massive CND, whose gas mass is about $10\\%$ of the dynamical mass inside $\\sim 0.1-1\\,{\\rm kpc}$. On the other hand, in a phase where the BH grows (i.e., a proto-QSO phase), the proto-QSO has a gas rich massive CNDs whose gas mass is comparable to the dynamical mass. Compared with the observed properties of the distant QSO SDSS J1148+5251 observed at $z=6.42$, we predict that SDSS J1148+5251 corresponds to the scenario of the late growth of SMBH with $z_{\\rm i}\\sim10$, which is accompanied by a massive CNDs with $M_{\\rm g}\\approx 5\\times 10^{10}M_{\\odot} $ and the luminous nuclear starburst $L_{\\rm SB}$ at infrared band with $L_{\\rm SB}\\approx 10^{47}\\,{\\rm erg}\\,{\\rm s}^{-1}$. Moreover, we predict that the progenitor of SDSS J1148+5251 can be the super-Eddington object. These predictions can be verified by ALMA, SPICA and JWST. ",
        "introduction": "Supermassive black holes (SMBHs) with masses in the range of $10^{6}-10^{9}M_{\\odot}$ are the engines that power active galactic nuclei (AGNs) and quasars (QSOs). There is also ample evidence that SMBHs reside at the center of most galaxies (e.g., Kormendy \\& Richstone 1995; Richstone et al. 1998; Ho et al. 1999), including the Milky Way (e.g., Genzel et al. 1997; Sch{\\\"o}del et al. 2002; Ghez et al. 2003). In the local Universe, there are tight correlations between the masses of SMBHs and the masses and velocity dispersions of the spheroidal components (bulge) of the hosts (e.g., Kormendy \\& Richstone 1995; Richstone et al. 1998; Ferrarese \\& Merritt 2000; Gebhardt et al. 2000; Tremaine et al. 2002; Marconi \\& Hunt 2003; H{\\\"a}ring \\& Rix 2004; Barth et al. 2005). These observational facts suggest that an intriguing link between the formation of bulges and SMBHs.  % High redshift QSOs are essential to understand the formation and evolution of SMBHs. In recent years, more than thirty QSOs at $z\\approx 6$ have been discovered (Fan et al. 2000, 2001, 2003, 2004, 2006; Goto 2006; Jiang et al. 2006; Willott et al. 2007; Jiang et al. 2009). The higher luminosity of the distant QSO at $z=6.42$, corresponding to SDSS J1148+3251 (Fan et al. 2003) with $\\sim 10^{47}\\,{\\rm erg}\\,{\\rm s}^{-1}$, implies that a SMBH with mass $\\geq 10^{9}M_{\\odot}$ is already in place within $\\sim$ 1 Gyr after the big bang, by assuming the Eddington luminosity. These requirements set significant constraints on the evolution and formation of SMBHs in the early Universe. There are some analytical and semi-analytical studies (Haiman \\& Loeb 2001; Haiman 2004; Yoo \\& Mirald-Escud\\'{e} 2004; Shapiro 2005; Volonteri \\& Rees 2006; Tanaka \\& Haiman 2009) and numerical simulations (Li et al. 2007; Sijacki et al. 2009) that discuss the formation of $> 10^{9}M_{\\odot}$ SMBH at $z >6$. % A key physical process governing evolution of QSOs is the mass accretion toward a SMBH, although BH growth via merging may be expected at the high mass BH with $M_{\\rm BH} > 10^{9}M_{\\odot}$ (e.g., Shankar 2009 and references therein). Although the accretion process is a complicated phenomenon over nine orders of magnitude in the size scale, the Eddington-limited accretion rate was simply assumed in most previous studies. However, the evolution of the SMBH is not only controlled by the accretion processes in the vicinity of the SMBH, but also related to the mass supply from the host galaxy to the circumnuclear region.   % A number of mechanisms for gas accumulation on the $\\sim 1-10$ kpc scale down to the galactic central region have been proposed, e.g., the tidal torque driven by the major and minor mergers of galaxies (e.g., Toomre \\& Toomre 1972; Mihos \\& Hernquist 1994, 1996; Saitoh \\& Wada 2004; Saitoh et al. 2009) and the stellar bars (e.g., Noguchi 1988; Shlosman et al. 1990; Fukuda et al. 1998; Fukuda et al. 2000; Maciejewski et al. 2002; Namekata et al. 2009), gas drag and dynamical friction in the dense stellar cluster (Norman \\& Scoville 1988) and the radiation drag (e.g., Umemura et al. 1997; Umemura 2001; Kawakatu \\& Umemura 2002). However, the accumulated gas does not directly accrete onto a SMBH, since the angular momentum of the gaseous matter cannot be thoroughly removed. Thus, some residual angular momentum would terminate the radial infall, so the accreted gas forms a reservoir, i.e., a circumnuclear disk (CND), in the central $\\sim$100 pc around a SMBH whose scale depends on the angular momentum of the gas. If the gravitational instability takes place in a CND, star formation is naturally expected in this region. Active nuclear star formation ($< 100\\,{\\rm pc}$) has been actually observed in nearby AGNs (e.g., Imanishi \\& Wada 2004; Davies et al. 2007; Watabe et al. 2008; Chen et al. 2009; Hicks et al. 2009). It is theorized that the nuclear starburst could obscure some types of AGNs (e.g., Ohsuga \\& Umemura 1999; Wada \\& Norman 2002: hereafter WN02; Thompson et al. 2005; Watabe et al. 2005; Ballantyne 2008; Schartmann et al. 2008; Wada et al. 2009). The nuclear starburst also affects the growth of SMBHs, because the radiation and/or supernova feedback from starbursts can enhance the mass accretion onto a SMBH (e.g., Norman \\& Scoville 1988; Umemura et al. 1997; WN02; Vollmer \\& Beckert 2003; Vollmer, Beckert \\& Davis 2008; Collin \\& Zahn 2008). In order to reveal the final rate of mass accretion to the BH region, it is crucial to {\\it link mass accretion processes from a galactic scale with those from an accretion disk in the vicinity of a central BH, via the CND}. Recently, we have proposed a theoretical model of a nuclear starburst disk supported by the turbulent pressure led by supernova explosions (Kawakatu \\& Wada 2008: hereafter KW08). In KW08, the turbulence excited by supernovae transports the angular momentum. We showed how a SMBH grows from a seed BH, taking into account the mutual connection between the mass-supply from a host galaxy and the physical states of the CND accompanied by the star formation. This theoretical model should be confirmed for the formation of high-$z$ ($z >6$) QSOs whose masses are $\\simeq 10^{9}M_{\\odot}$.   This paper is organized as follows. We briefly outline the  model in KW08 in $\\S 2$. By adopting KW08, we will show the necessary conditions to form QSOs at $z > 6$ in terms of the total accreted gas mass from host galaxies and star formation efficiency in the CNDs ($\\S 3$). In addition, we will predict observable properties and the evolution of QSOs at $z >6$ ($\\S4$). In section 5, we discuss how the distant QSO J1148+5251 at $z=6.42$ forms and predict the nature of the early phase of this QSO. Section 6 is devoted to our summary. Unless otherwise stated, all results shown below refer to the currently favored $\\Lambda$ cold dark matter model with $\\Omega_{\\rm M}=0.24$, $\\Omega_{\\Lambda}=0.76$, $h=0.73$, $\\Omega_{\\rm b}=0.042$, $\\sigma_{8}=0.74$ and $n=0.95$ (Spergel et al. 2007). ",
        "conclusions": "Based on the physical model of supermassive black hole (SMBH) growth via gas accretion in $\\sim $100 pc scale circumnuclear disks (CNDs) proposed by Kawakatu and Wada (2008), we investigate the formation of high-$z$ ($z > 6$) QSOs whose BH masses are $ M_{\\rm BH}> 10^{9}M_{\\odot}$. We show the necessary conditions to form QSOs at $z > 6$, in terms of the total gas mass accreted  from host galaxies, $M_{\\rm sup}$ and star formation efficiency in the CND, $C_{*}=\\dot{M}_{*}/M_{\\rm g}$ where $\\dot{M}_{*}$ and $M_{\\rm g}$ are the star formation rate and the gas mass, respectively. Our main conclusions are summarized as follows.  \\begin{enumerate} \\item The required conditions for the formation of QSOs at $z >6$ are as follows; (i) Rapid and large mass supply with $M_{\\rm sup} > 10^{10}M_{\\odot}$ within 1 Gyr from hosts to CNDs, because the final BH mass is only $1-10\\%$ of the total mass supplied from hosts. The fraction is determined by the efficiency with which SN energy is transferred to the cold gas, and the efficiency of angular momentum transfer due to turbulent viscosity. (ii) High star formation efficiency, i.e., $10^{-8}\\,{\\rm yr}^{-1}< C_{*} < 10^{-6}\\,{\\rm yr}^{-1}$ is comparable to high-$z$ starburst galaxies such as SMGs. (iii) Rapid mass inflow from hosts with a period of the mass-supply, $t_{\\rm sup}< 5\\times 10^{9}\\,{\\rm yr}$.   \\item We find that if the BH growth is limited by the Eddington accretion, the final BH mass is greatly suppressed when the period of mass-supply from hosts, $t_{\\rm sup}$, is shorter than the Eddington timescale. Thus, the super-Eddington growth is required for the QSO formation, while $t_{\\rm sup}$, which is determined by the efficiency of angular momentum transfer, is shorter than $\\sim 10^{8}\\,{\\rm yr}$.  \\item On the basis of these results, we predict the observable properties and the evolution of QSOs at $z >6$ as follows; (i) The evolution of the QSO luminosity strongly depends on the redshift $z_{\\rm i}$ at which the accretion onto a seed BH got started. In other words, the bright QSOs given $z$ (e.g., $z=6$) favor the late growth of SMBHs (i.e., $z_{\\rm i}\\approx 10$) rather than the early growth (i.e., $z_{\\rm i}\\approx 30$) because the timescale of gas supplied from hosts is shorter than that of the star formation in the CND. (ii) In the QSO phase, after which the gas supply stopped and when the AGN luminosity is above the threshold one, $L_{\\rm AGN, th}$ (e.g., $L_{\\rm AGN, th}=10^{46}\\,{\\rm erg}\\, {\\rm s}^{-1}$), there exists a stellar rich massive CND with $M_{\\rm g}/M_{\\rm dyn} \\approx 0.1$ where $M_{\\rm dyn}$ is the dynamical mass of the CND plus BH system. On the other hand, in the proto-QSO phase (the early phase of a growing BH), the proto-QSO has a gas rich massive CNDs with $M_{\\rm g}/M_{\\rm dyn}\\approx 1$.  \\item Our theoretical model predicts that the distant QSO SDSS J1148+5251 discovered at $z=6.42$ favors the model of late growth of SMBH with $z_{\\rm i}=8$ rather than early growth with $z_{\\rm i}=25$. If this is the case, the QSO  harbors a massive CNDs with $M_{\\rm g}= 5\\times 10^{10}M_{\\odot}$ and the luminous nuclear starburst in the infrared band with $L_{\\rm SB}\\approx 10^{47}\\,{\\rm erg}\\,{\\rm s}^{-1}$. These predictions can be checked by the future observations (e.g., ALMA, SPICA, and JWST), in order to judge whether our model is reasonable for the formation of SMBHs. In addition, we predict that SDSS J1148+5251 have been evolved in a massive dark matter halo and/or a massive host (i.e., $M_{\\rm DM}\\approx 10^{12-13}M_{\\odot}$ and/or $M_{\\rm host}\\approx 10^{13}M_{\\odot}$), in order to build a SMBH with $>10^{9}M_{\\odot}$ only by gas accretion. The progenitor of SDSS J1148+5251 can be super-Eddington objects and observed as the high-$z$ ULIRGs such as SMGs with the infrared luminosity $L_{\\rm IR}\\sim 10^{47}\\,{\\rm erg}\\,{\\rm s}^{-1}$. The progenitor has a {\\it gas rich} CND with the mass ratio $M_{\\rm g}/M_{\\rm dyn}\\approx 1$ and $M_{\\rm g}\\geq 10^{9}M_{\\odot}$, which will be observed by ALMA.  \\end{enumerate}"
    },
    "0910/0910.4571_arXiv.txt": {
        "abstract": "We investigate data compression schemes for proposed all-sky diffraction-limited visible/NIR sky surveys aimed at the dark energy problem.  We show that lossy square-root compression to 1 bit of noise per pixel, followed by standard lossless compression algorithms, reduces the images to 2.5--4 bits per pixel, depending primarily upon the level of cosmic-ray contamination of the images.  Compression to this level adds noise equivalent to $\\le 10\\%$ penalty in observing time.  We derive an analytic correction to flux biases inherent to the square-root compression scheme.  Numerical tests on simple galaxy models confirm that galaxy fluxes and shapes are measured with systematic biases $\\lesssim10^{-4}$ induced by the compression scheme, well below the requirements of supernova and weak gravitational lensing dark-energy experiments.  An accompanying paper \\citep{Ali} bounds the shape biases using realistic simulated images of the high-Galactic-latitude sky.  The square-root preprocessing step has advantages over simple (linear) decimation when there are many bright objects or cosmic rays in the field, or when the background level will vary. ",
        "introduction": "The unexpected discovery of the acceleration of the Hubble expansion of the Universe has inspired numerous proposals for large-scale observational projects to seek further constraints on this ``dark energy'' problem \\citep[see][for a review of dark energy]{FTH}.  Several such proposals involve space-based visible or NIR imaging surveys of a substantial fraction of the full celestial sphere at resolution near the diffraction limit of 1--2~meter telescopes.  Such surveys would also be extraordinarily valuable for general astrophysics, but the amount of data involved is large.  The number of pixels in a survey of a fraction $f_{\\rm sky}$ of the sky, with pixel scale $p$ and $N_{\\rm exp}$ exposures per sky location, will be \\begin{equation} N_{\\rm pix} = 2.7\\times 10^{14} \\left({ p \\over 0\\farcs1 }\\right)^{-2} { f_{\\rm sky} N_{\\rm exp} \\over 5}. \\end{equation} A crate filled with disk drives is all that is needed these days for transmission of hundreds of terabytes of data between terrestrial sites.  The transmission of this data from a spacecraft at the Earth-Sun L2 Lagrange point 1.5~million km away is, however, a substantial expense.  If the data require an average of bpp bits per pixel to transmit, with a telemetry rate $r$, and a survey duration of $T$, then the number of hours of downlink per day required will be \\begin{equation} \\mbox{downlink time} ={\\rm bpp}\\times 0.33 {\\rm hrs \\over day} { f_{\\rm sky} N_{\\rm exp} /T \\over 1.25/{\\rm yr}} \\left({r \\over 150\\,{\\rm Mbps}}\\right)^{-1} \\left({ p \\over 0\\farcs1 }\\right)^{-2}. \\end{equation} To avoid excessive resource use both on the spacecraft and ground stations, the data volume must be reduced well below the 16 bits per pixel that is typical for digitization of astronomical array detectors. In this paper we investigate means for compressing astronomical imaging data well below this, ideally ${\\rm bpp} \\lesssim 4$.  Significant compression of the detector output signals will require a lossy compression step to reduce the bit depth of the noise---sometimes called ``pre-processing''---followed by a lossless algorithm to reduce the mean bit depth to the signal entropy level.  Because most of the pixels in an extragalactic visible/NIR exposure will be close to the sky level, it is clear that the primary determinant of the final bpp value will be the number of noise bits retained by the lossy step in pixels near sky level---extensive demonstration of this for ground-based images is done by \\citet{PSW}. Optical astronomers are averse to lossy compression for fear of losing the information from their dearly acquired photons or fear of inducing biases in high-precision measurements.  It should be pointed out, however, that lossy compression is already being performed on essentially all data when the analog detector outputs are digitized.  Analog-to-digital gains are usually set, however, so that the errors induced by digitization are below the read noise of the detector system.  In radio astronomy it is well known that digitization to only 1 or 2 bits per sample can capture nearly all the signal information when $S/N$ per sample is low \\citep{Weinreb,TMS}.  The critical parameter for information loss is the number of bits of noise retained, {\\it i.e.} the ratio $b=\\sigma_N/\\Delta$ of the signal noise to the digitization step size. We seek a lossy compression scheme which holds $b$ fixed over the full dynamic range of the detector.  When the signal noise is dominated by shot noise from the detected photoelectrons, it is straightforward to show that a {\\em square-root compression scheme} holds $b$ constant.  The application of square-root compression to Poisson-dominated data has been discussed by \\citet{GowenSmith} and is part of the signal chain for the {\\it James Webb Space Telescope} \\citep{JWST} and {\\it Supernova Acceleration Probe} \\citep{LinMarriner} spacecraft designs.  Further discussion is in \\citet{SWP}. In \\S\\ref{algo} we review the algorithms of such a scheme and \\S\\ref{noisesec} shows that the square-root compression algorithm increases the noise by a fraction $1/24b^2$ for cases of interest.  Once a lossy compression algorithm is determined to reduce the data volume with acceptably small degradation of the noise level, it is critical to determine whether the codec process will induce any bias on the astronomical measurements from the data.  Future dark energy experiments will require very demanding, high-precision measurements.  For use of Type Ia supernovae as distance indicators, systematic errors in flux measurements must be $\\ll 1\\%$.  For determination of galaxy redshifts using the ``photo-z'' method, similarly small systematic errors must be present in galaxy photometry.  We target codec-induced systematic errors in flux to be $<10^{-3}$ so that they contribute negligibly to the photometric error budget.  Even more demanding are programs for weak gravitational lensing (WL) measurements, which depend upon determination of the shapes (ellipticities) of $\\gtrsim 10^9$ galaxies on the sky.  \\citet{HTBJ} and \\citet{AmaraRefregier} both conclude that measurements of galaxy ellipticities must have systematic errors of $\\le 10^{-3}$ in order to remain subdominant to statistical errors in the WL experiments. Because there are many possible sources of shape measurement error---primarily uncorrected instrumental image distortions---we would like to have shape biases induced by the codec be a minor contributor to the shape error budget, {\\it i.e.} the shift $\\delta e_i$ in each component $e_i$ of galaxy ellipticity should on average be $<10^{-4} e_i$.  The goal of this paper is to examine the biases induced by square-root compression down to $b=1$ bit of noise per pixel, and to show how these biases can be reduced to negligible levels.  A companion paper \\citep{Ali} tests for codec-induced shape bias on realistic simulations of space-based CCD imaging.  Many lossy image compression algorithms are based on decomposition of 2d images over a function library, followed by truncation of the terms in the library---for example wavelet transforms \\citep{White,Press} or the JPEG algorithms.  Even more radical compression is possible with ``compressed sensing'' techniques if the data are known to be sparse in some library \\citep{Candes}.  For WL applications we are wary of lossy compression algorithms that examine groups of pixels and might hence induce correlated codec errors over multiple pixels.  We therefore focus here on a lossy step that compresses data strictly on a pixel-by-pixel basis.  If single-pixel codec errors are unbiased and completely uncorrelated between pixels,  object shape measurements should be unbiased.  We investigate the single-pixel codec in this work,  leaving open the question of whether grouped-pixel algorithms can meet the shape bias requirements of weak lensing surveys while obtaining better compression than single-pixel codecs.  In \\S\\ref{biassec} we show that there is a significant bias in the mean values reconstructed by the square-root class of codecs.  We also show, however, that a straightforward adjustment of the decoded values can remove this bias to remarkable accuracy, largely independent of the details of the probability distribution of the raw data.  We therefore expect the codec to produce unbiased flux and shape measurements, which we test with simplified galaxy images in \\S\\ref{testsec}.  \\citet{Ali} test for shape-measurement bias on data that is a realistic simulation of images expected from the {\\it Joint Dark Energy Mission (JDEM)}\\footnote{\\tt http://jdem.gsfc.nasa.gov} CCD imager.  Finally in \\S\\ref{bppsec} we compress realistic images with the lossy square-root algorithm followed by standard lossless algorithms, and measure the number of bits per pixel required to transmit the fully compressed images.  In the conclusion we summarize our findings, namely that it is possible to compress {\\it JDEM}-like images to 3--4 bits per pixel with only 4\\% penalty in image noise, and to reconstruct them such that biases in flux and shape measurements induced by the codec are not detected, down to $\\approx 10^{-4}$ levels. ",
        "conclusions": "Future dark-energy surveys such as {\\em JDEM} or {\\em Euclid}\\footnote{\\tt http://sci.esa.int/euclid} propose to survey a substantial fraction of the full sky with space-based visible/NIR imaging near the diffraction limit.  A data compression scheme is needed to satisfy telemetry constraints, but the weak gravitational lensing measurements place very strict requirements on the degree to which the compression algorithm can bias the measured shapes or fluxes of galaxies.  For this reason we choose to investigate lossy compression processes that operate on a pixel-by-pixel basis, fearing than algorithms examining pixel groups might induce correlated codec errors between pixels which bias the galaxy ellipticities.  Square-root pre-processing is known to be a good means for compressing shot-noise-dominated data streams while keeping a fixed number $b$ of bits per $\\sigma$ of noise across the full detector dynamic range.  The reduction of entropy by reducing the number of noise bits is essential to efficient subsequent lossless compression algorithms.  We derive an expression for bias in the naive reconstruction of square-root-compressed data and present a correction scheme, which we numerically verify is successful to high accuracy.  Since the noise induced by roundoff in the compression process increases the required observing time by $1/12b^2$, we confine our examination to $b\\ge1$.  We verify that flux and shape measurements of simple galaxy models are biased by $\\lesssim 10^{-4}$ for a range of galaxy sizes, ellipticities, orientations, and signal-to-noise ratios.  This is well within the weak lensing requirements.  We verify this with simple postage-stamp galaxy images; \\citet{Ali} bound the shape biases determined from realistic simulations of the true extragalactic sky.  At $b=1$ the Shannon entropy of the square-root-compressed images varies from 2.1 bits per pixel for pure noise to 3.6 bits per pixel for images with the expected astronomical scene and triple the expected number of cosmic rays.  Lossless compression algorithms such as {\\tt bzip2} realize compress levels within 0.3--0.4 bits of the Shannon limit.   For simulated images best resembling expected {\\em JDEM} CCD data, we find compression to 3.1 bits per pixel using {\\tt bzip2}.  The size of the transmitted data stream depends primarily on the degree of cosmic-ray contamination once $b=1$ is set: schemes to more efficiently compress or mask cosmic rays before compression might lead to lower bit rates.  Higher cosmic-ray rates or more crowded astronomical fields, {\\it e.g.} near the Galactic plane, will require more bits.  Since pure noise images require 2.4 bits per pixel with {\\tt bzip2} at $b=1$, we can state that compression to $<2.4$ bits per pixel will require $b<1$, and hence higher observing-time penalty.  We thus validate square-root lossy compression as useful for {\\em JDEM} data and capable of meeting requirements for low bias.  Other approaches would be feasible as well. For instance simple decimation to a fixed code step $\\Delta$ that equals the sky noise level should produce similar noise and bias properties.  The presence of cosmic rays and/or bright objects in $\\gtrsim 10\\%$ of the pixels gives the square-root codec a useful advantage in mean number of required telemetry bits per pixel.  The square-root algorithm also has a robust and exceptionally simple implementation---a single lookup table---meaning that there is no need for on-board processing to recognize cosmic rays or determine the sky noise level.  It is thus expected that visible/NIR dark-energy extragalactic imaging survey data can be transmitted to the ground with 2.5--4 bits per pixel, depending upon the cosmic-ray flux level and the sophistication of onboard cosmic-ray identification or removal.  This compression can be achieved with no significant bias of observed quantities, and with 8\\% penalty in observing time to achieve fixed signal-to-noise ratio."
    },
    "0910/0910.3398_arXiv.txt": {
        "abstract": "The Fermi Gamma-Ray Space Telescope was launched in June 2008 and the onboard Large Area Telescope (LAT) has been collecting data since August of that same year. The LAT is currently being used to study a wide range of science topics in high-energy astrophysics, one of which is the study of high-energy cosmic rays. The LAT has recently demonstrated its ability to measure cosmic-ray electrons, and the Fermi LAT Collaboration has published a measurement of the high-energy cosmic-ray electron spectrum in the 20 GeV to 1 TeV energy range. This talk will discuss the prospects for using the LAT to perform a similar analysis to measure cosmic-ray proton events. The instrument response for cosmic-ray protons will be characterized and an assessment of the potential to measure the cosmic-ray proton energy spectrum will be presented. ",
        "introduction": "The Fermi Gamma-ray Space Telescope was designed for the study of many interesting topics in astrophysics ranging from pulsars, active galactic nuclei, gamma-ray bursts, and the indirect detection of dark matter.  Another intriguing subject that has attracted strong interest is the use of the Fermi LAT in the study of cosmic rays.  Early in the design stages of the LAT, the potential for making a cosmic-ray electron measurement was acknowledged \\cite{Ormes}, \\cite{Moiseev}. Recently the Fermi LAT Collaboration demonstrated this ability and published a high-statistics measurement of the cosmic-ray electron spectrum in the energy range from 20 GeV to 1 TeV \\cite{CREpaper}, containing about 4.5 million events collected over a six month period from August 2008 to January 2009.  The electron spectrum measurement can be used in the study of cosmic-ray propagation and in the constraint of the diffuse gamma-ray emission \\cite{CRPropPaper}.  Furthermore, it may also be possible to use the LAT to obtain a measurement of the cosmic-ray proton spectrum, and studies are being conducted within the LAT collaboration to explore this possibility.  We present some preliminary findings illustrating the prospects for such an analysis.  First a general introduction to the Fermi LAT instrument will be given, followed by a brief description of the analysis performed for the electron spectrum measurement.  Then we will present a discussion of the prospects for a cosmic-ray proton spectrum analysis. ",
        "conclusions": ""
    },
    "0910/0910.3021_arXiv.txt": {
        "abstract": "We present new {\\it Chandra} observations of the radio galaxy 3C 445, centered on its southern radio hotspot.  Our observations detect  X-ray emission displaced upstream and to the west of the radio-optical hotspot. Attempting to reproduce both the observed spectral energy distribution (SED) and the displacement, excludes all one zone models. Modeling of the radio-optical hotspot spectrum suggests that the electron distribution has a low energy cutoff or break approximately at the proton rest mass energy.  The X-rays could be due to external Compton scattering of the cosmic microwave background (EC/CMB) coming from the fast (Lorentz factor $\\Gamma\\approx 4$) part of a decelerating flow, but this requires a small angle between the jet velocity and the observer's line of sight ($\\theta\\approx 14^{\\circ}$). Alternatively, the X-ray emission can be synchrotron from a separate population of electrons.  This last interpretation does not  require the X-ray emission to be beamed. ",
        "introduction": "}   The hotspots of powerful Fanaroff-Rilley type II [FRII;  Fanaroff \\& Riley 1974] radio galaxies and quasars are the locations where the jets of these sources, after propagating for distances up to $\\sim 1$ Mpc, terminate in a collision with the intergalactic medium. The optical emission observed in several hotspots  suggests  that at least a fraction of the electrons that goes through the shock(s) formed at the termination  of the jets undergoes efficient particle acceleration \\citep{heavens87, meisenheimer89, prieto02}.  {\\sl Chandra} results show that while in some sources (e.g. Cygnus A; Wilson et al. 2000) the hotspot X-ray emission is consistent with synchrotron-self Compton radiation from relativistic electrons  in energy equipartition with the magnetic field (SSCE), in other sources (e.g. the hotspot on the jet side of Pictor A; Wilson et al. 2001,  Hardcastle \\& Croston 2005, Migliori et al. 2007, Tingay et al. 2008) the X-ray emission is at a much higher level  (by up to a factor of $ \\sim 1000$). The nature of this anomalously bright (significantly brighter than SSCE) X-ray hotspot emission  remains a matter of active discussion in the literature, being connected to the issue of particle acceleration efficiency  and jet power [for two recent  reviews see Harris \\& Krawczynski (2006) and Worall (2009)] and extending to a similar and possibly related issue for the knots of powerful jets (e.g. Kataoka \\& Stawarz 2005).  Two different considerations appear to be relevant to the collective properties of hotspot X-ray emission, relativistic beaming and hotspot  luminosity; however, it is not as yet clear how exactly  these may be combined to reproduce the  observed phenomenology. Based on early {\\sl Chandra} results, suggesting that in most cases the anomalously bright X-ray hotspots were seen at the approaching  jet of  sources with jets forming relatively small angles to the observer's line of sight, Georganopoulos \\& Kazanas (2003) proposed that  the X-ray emission is  beamed, and that the plasma in the hotspot is relativistic and decelerating from a bulk Lorentz factor $\\Gamma\\sim 2-3$ to velocities that match the subrelativistic advance speed of lobes  ($u/c\\approx 0.1$; Arshakian \\& Longair 2001). In this scenario the X-ray emission  is mostly due to upstream Compton (UC) scattering of electrons in the upstream fast ($\\Gamma\\sim 2-3$) part of the flow,  inverse-Compton scattering the  synchrotron photons produced downstream. The relevance of beaming is also supported by a study of  quasar hotspots by Tavecchio et al. (2005) that concluded that the anomalously bright X-ray emission is indeed found mostly on the hotspot corresponding to the approaching jet side, although these authors favor the cosmic microwave background as a source of seed photons. On the other hand, Hardcastle et al. (2004) find only weak evidence that the anomalously X-ray bright hotspots are more frequently found at the termination of the approaching jet.  The second consideration, namely the hotspot luminosity, was introduced by \\cite{brunetti03}, who found that the synchrotron emission seen at radio energies extends to the optical for lower power sources like 3C 445 \\citep{prieto02}  but cuts off before the optical regime for powerful sources like Cygnus A \\citep{wilson00}.  They explained this as a consequence of radiative losses increasing with hotspot luminosity. The luminosity range over which this decrease of the synchrotron peak frequency with increasing luminosity is observed, was extended to powerful hotspots by \\cite{cheung05}. Based on an extensive sample of hotspot multiwavelength data, \\cite{hardcastle04} argued for the relevance of the hotspot luminosity to the X-ray emission, by showing that  hotspots with X-ray luminosity much higher than that predicted by SSCE % were usually of  low luminosity, in contrast to more powerful hotspots.  % However, instead of UC, \\cite{hardcastle04} favored synchrotron emission from an altogether separate electron population as the source of the anomalously bright X-ray emission. Other workers have also come to the conclusion that synchrotron radiation from a second electron population is the most likely emission mechanism for this component (e.g., in the case of Pic A, see Fan et al. 2008, Tingay et al. 2008), partly as a result of observing compact radio hotspots with the VLBA.  Synchrotron radiation, but  from a single electron population, is indeed the X-ray emission mechanism for the jets of low power FR I radio galaxies, as seen from their single component radio-optical-X-ray spectra (e.g. Perlman \\& Wilson 2005 for the jet of M 87). This turns out to be the case also for some of the weakest hostspots of FR II radio galaxies: the northern hotspot of 3C 390.3 \\citep{hardcastle07}, the northern hotspots of 3C 33 \\citep{kraft07}, the eastern hotspots of 3C 403 \\citep{kraft07} and both hotspots of $0836+299$ \\citep{tavecchio05}. It comes as no surprise that in these sources  their X-ray emission is much higher than the anticipated SSC flux, since the X-ray emission is indeed the continuation of the radio-optical synchrotron spectrum to higher energies. We therefore do not consider sources that exhibit a single spectral component from radio to X-ray emission to be part of the family  of hotspots exhibiting anomalously high X-ray emission.  An additional handle in understanding the X-ray emission process of the anomalously X-ray bright hotspots (those for which the X-rays (i) cannot be a continuation of the synchrotron spectrum and (ii) are significantly brighter than predicted by SSCE) can come from spatial displacements between the emission at different frequencies. In a handful of these hotspots, displacements between the radio and the X-ray hotspot emission have been observed, with the X-rays being upstream of the radio: 3C 351 \\citep{hardcastle02}, 4C 74.26 \\citep{erlund07}, 3C 390.3 and  3C 227 \\citep{hardcastle07}, 3C 321 \\citep{evans08}, 3C 353 \\citep{kataoka08}. {\\sl So far no displacements have been observed in the hotspots of sources for which their X-ray emission is consistent with SSCE.} An additional important characteristic is that when optical-IR emission is detected from these hotspots (as in 3C227, 3C 390.3, 3C 351) it coincides with or is shifted somewhat upstream of the radio, to a location downstream of the X-rays. Beaming in the hotspots that exhibit displacements can be a relevant but not a dominant influence, because, while in three of them the hotspots are in the approaching jet that points toward us (4C  74.26, 3C 227, 3C 351),   in two others (3C 321 and 3C 353), the jets are believed to be close to the plane of the sky, while in superluminal 3C 390.3 the  hotspot exhibiting the X-ray radio displacement is  at the termination of the counter jet. Similarly, a low hotspot power does not seem to be strictly required, since the hotspot of 3C 351 is rather powerful and still exhibits the displacements mentioned above.  In this paper we present {\\sl Chandra} X-ray observations and discuss the X-ray emission mechanism of the hotspots of 3C 445, a broad line FR II radio galaxy at redshift $z=0.0562$  \\citep{eracleous94}, which for the standard cosmology ($H_0=71 $ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_\\Lambda=0.73$ and $\\Omega_M=0.27$) corresponds to a luminosity distance of $247.7 $ Mpc. This is a promising source for constraining the hotspot X-ray emission mechanism with high resolution {\\sl Chandra} observations.  Both its southern and northern  hotspots have been detected in the radio and in the near-IR/optical, with the southern hotspot being $\\sim 3$ times brighter than the northern and having multiple sites of synchrotron optical emission, a manifestation of ongoing particle acceleration \\citep{prieto02,mack09}.  The hotspot - counter hotspot luminosity difference can be intrinsic, or due to mild beaming with the southern hotspot being on the approaching jet side.  In the first case, adopting the suggestion of \\cite{hardcastle04}, we expect X-ray emission brighter than SSCE from both hotspots, because this is a source with a modest extended power at  $178$ MHz, $P_{178}=3.0 \\times 10^{25}$ W Hz$^{-1}$ sr$^{-1}$  \\citep{hardcastle98}, as well as a low $5$ GHz  luminosity from both hotspots  ($L_{5 GHz}=3.4 \\times 10^{22}$ W Hz$^{-1}$ sr$^{-1}$ for the northern hotspot and $L_{5 GHz}=7.9 \\times 10^{22}$ W Hz$^{-1}$ sr$^{-1}$ for the southern hotspot \\citep{mack09}). In this scenario, the X-ray emission from the weaker northern hotspot is expected to be higher than predicted by the X-ray to radio ratio from the brighter southern hotspot.  This model makes no prediction regarding offsets between emission in different bands.  In the second case, we expect that due to beaming, the presumed approaching-jet side, southern hotspot will have more pronounced X-ray emission.  In the context of a relativistic decelerating hot spot flow  \\citep{georganopoulos03, georganopoulos04},  we also expect to observe  offsets, with the X-rays peaking upstream of the radio. If, in addition, distributed particle acceleration takes place (as suggested by Prieto et al. 2002), the optical and radio emission will peak at the same location (see Figure 3 of Georganopoulos \\& Kazanas 2004).   3C 445 has been the subject of several X-ray observations.  {\\it GINGA} observations were reported by Pounds (1990), while {\\it ASCA} observations were reported by Sambruna et al. (1998).  Both observations reported an absorbed Seyfert nucleus with a hard X-ray continuum. {\\it XMM-Newton} observations of 3C 445 were previously published by \\cite{sambruna07} and  \\cite{grandi07},  but in those data a partial window configuration was chosen for the MOS and pn that caused both the northern and southern hotspots to be at locations where the CCDs were not read out. This work thus represents the first study of extended X-ray emission from this object, and we report the discovery of X-ray emission from  the southern hotspot, as well as a meaningful X-ray flux upper limit from the northern hotspot. In \\S 2 we discuss the {\\it Chandra} observations, as well as the data reduction and analysis procedures.  In \\S 3 we present and contrast the morphology seen in each band for the southern hotspot. In \\S 4 we  discuss the  X-ray emission mechanism of the southern hotspot, present our conclusions and suggest directions for future work. ",
        "conclusions": "}  One of the key results of our {\\sl Chandra}  observations is  that the X-ray emission of the southern hotspot has a very different morphology than that seen in the radio and optical and also shows a likely displacement. {\\sl This displacement rules out all forms of one-zone models.}  Before discussing the possible interpretations for the X-ray emission, we turn to the radio and optical emission that are   approximately  cospatial. The radio-optical SED of the southern and northern hotspots are  plotted in Figure \\ref{fig:sed}, along with the X-ray points. The southern hotspot is brighter in the radio and optical by a factor of $\\approx 3$. If we attribute the brightness difference to beaming, we are forced to conclude that the beaming of the radio-optical emitting plasma  is mild. Below, we discuss the multiwavelength emissions of both hotspots and model possible X-ray emission mechanisms.  \\subsection{A  high value of $\\gamma_{min}$ in the radio-optical hotspot?}  The southern hotspot's radio--optical SED can be modeled as synchrotron emission from a population of relativistic electrons in energy equipartition with the hotspot magnetic field.  Assuming that beaming is not important for the radio--optical emission, the equipartition magnetic field is \\begin{equation} \\displaystyle B_{eq}=\\left[{96 \\pi^2  m_e c  L_{r} \\nu_r^{(s-1)/2} (\\gamma_{min}^{2-s}-\\gamma_{max}^{2-s})\\over c_1^{(s-3)/2} \\sigma_\\tau  V (s-2)}\\right]^{2/(s+5)}, \\label{eq:equip} \\end{equation} where $L_r$ is the radio luminosity at frequency $\\nu_r$, $m_e$ is the electron mass, $c$ is the speed of light, $\\sigma_\\tau$ is the Thomson cross section, and $V$ is the volume of the emitting region. The injected electron distribution is a power law of index $s=2\\alpha_r+1=2.6$ from Lorentz factor $\\gamma_{min}$ to $\\gamma_{max}$, with  $\\alpha_{ro}=0.9=\\alpha_r$ being the radio-optical spectral index \\citep{mack09}. To derive  equation (\\ref{eq:equip}), we  assumed that an electron  of Lorentz factor $\\gamma$ in a magnetic field $B$ radiates most of its energy at the characteristic frequency $\\nu=c_1 B \\gamma^2$, with $c_1=e/(2\\pi m_e c)$. In equation (\\ref{eq:equip}), $\\gamma_{max}$ can in many cases be determined observationally from the maximum observed synchrotron frequency. However, $\\gamma_{min}$ is customarily set to a value chosen by hand. As we discuss now, there is a way to  determine observationally, or at least constrain the value of $\\gamma_{min}$, and through this get a more appropriate value for $B_{eq}$. This in turn affects significantly the level of both the SSC and the EC/CMB emission.  Because in this case $s>2$ and the synchrotron emission extends for at least six decades in frequency, $\\gamma_{max}/\\gamma_{min}\\gg 1$, to a very good approximation $B_{eq}\\propto \\gamma_{min}^{-2(s-2)/(s+5)}=\\gamma_{min}^{-1.2/7.6}$: an increase in $\\gamma_{min}$ results in a mild decrease of $B_{eq}$. There are two observational constraints on $\\gamma_{min}$. An upper limit on $\\gamma_{min}$ is derived from the fact that it has to be low enough to produce the lowest observed radio frequency from the hotspot $\\nu_{r, min}  > c_1 B_{eq} \\gamma_{min}^2$.  A lower limit comes from the fact that there is no sign of radiative cooling in the radio-optical SED, because the optical flux level is found practically on the extrapolation of the   radio spectrum. This sets an upper limit on the magnetic field in the radio-optical hotspot (regardless of equipartition arguments), which in turn sets a lower limit on $\\gamma_{min}$.  To demonstrate these considerations, in Figure \\ref{fig:sed} we plot with a thin solid line the synchrotron emission in the case of $\\gamma_{min}=1$. This corresponds to  an equipartition magnetic field, $B_{eq}=70.4\\; \\mu$G. As  can be seen, while the synchrotron SED clearly extends below the lowest radio frequency safely associated with the hotspot [this is the 4.8 GHz point,  because the 1.4 GHz point may be contaminated with non-hotspot emission, (see \\cite{mack09, prieto02}) as is also true for the lower frequency points (Kassim et al. 2007).], the synchrotron spectrum breaks at $\\sim 10^{12-13}$ Hz and by doing so, underproduces the optical emission of the hotspot. This is because the high value of $B_{eq}$ causes a break in the electron energy distribution due to radiative cooling (a cooling break is expected at $\\gamma_b=3m_e c^2 /[4 \\sigma_\\tau (B^2/8\\pi+U_{CMB}) R]$, where $U_{CMB}$ is the energy density of the cosmic microwave background and $R$ is the size of the  hotspot that determines the electron escape time $R/c$). To fit the observed SED, we need significantly higher values of $\\gamma_{min}$. A value of $\\gamma_{min}\\approx 1840$ similar to the proton to electron mass ratio $m_p/m_e$ is required to ensure that  there is no cooling break signature at frequencies lower than optical. The value of the corresponding equipartition magnetic field is  $B_{eq}=21.5 \\; \\mu$G. We plot the resulting synchrotron SED in Figure  \\ref{fig:sed} with a thick solid line. Note that at low frequencies the model with $\\gamma_{min}\\approx 1840$ (thick solid line)  exhibits a break due to the high value of $\\gamma_{min}$ (the slope below the break is due to the $\\nu^{1/3}$ synchrotron emissivity of electrons with Lorentz factor $\\gamma_{min}$). Note also that this model manages to  reproduce the optical emission of the hotspot.  \\begin{figure} \\epsscale{1.10} \\plotone{3C445gmin1841.eps} \\caption{The SED of the southern  hotspot of 3C 445 is shown with  diamonds for the radio and optical and bow-tie for the X-rays. The SED of the northern hotspot is also plotted with asterisks, including the  upper limit for the X-ray flux. The data used are taken from Table 1. Due to angular resolution constraints, the three lowest radio frequencies may include lobe emission and should be considered upper limits for the hotspot fluxes. The thin (thick) lines represent emission in equipartition conditions, assuming $\\gamma_{min}=1$ ($\\gamma_{min}=1840$). Solid lines represent the synchrotron, short dash lines the SSC and long dash lines the EC/CBM emission.} \\label{fig:sed} \\end{figure}  There is practically little freedom for $\\gamma_{min}$ around $m_p/m_e$  if we want to model the radio to optical SED with synchrotron in equipartition. Let us mention that observationally driven arguments for similarly high  values of $\\gamma_{min}$ in hotspots of other radio galaxies have been presented by other astronomers (e.g. Blundell et al. 2006, Stawarz et al. 2007, Godfrey et al. 2009).  However, some jet sources require much lower values, e.g., PKS 0637--752 (Mehta et al. 2009). Values of $\\gamma_{min}\\approx m_p/m_e$ are particularly interesting, because this is the minimum energy that electrons crossing a shock must have to be picked up efficiently by Fermi acceleration (e.g. Spitkovsky 2008). The level of both the EC and SSC emission depend on the value of $B_{eq}$, which in turns depends on $\\gamma_{min}$:  because $L_{SSC, EC}/L_S \\propto U_B^{-1}\\propto B^{-2}$, and $B_{eq}\\propto \\gamma_{min}^{-2(s-2)/(s+5)}$, for $B=B_{eq}$, $L_{SSC, EC} \\propto \\gamma_{min}^{4(s-2)/(s+5)}=\\gamma_{min}^{2.4/7.6}$. Therefore, an  increase from $\\gamma_{min}=1$ to   $\\gamma_{min}=1840$, should increase the EC and SSC level by a factor of $\\approx 10.7$, as seen in Figure \\ref{fig:sed}. Even for $\\gamma_{min}=1840$,  the X-ray emission of the radio-optical hotspot is much weaker than the upstream detected  X-ray emission.   If analyses like the above point toward a high value of $\\gamma_{min}$, then future low frequency-high angular resolution observations ({\\it e.g.}, with the LWA) should detect the SED break at low frequencies. Existing low fequency radio observations suggest such a break for the eastern hotspot of Cygnus A \\citep{lazio06}. For 3C 445,  VLA observations at 74 MHz and 330 MHz (Kassim et al. 2007),  provide us with upper limits for the  emission of the hotspots (due to lobe contamination). The low frequency data plotted in figure \\ref{fig:sed} may be contaminated by lobe emissions and hence are not sufficient to constrain the spectral shape below 4.8 GHz. If a  low frequency break is found by future low frequency-high resolution observations, they will strengthen the above picture. If on the other hand the low frequency radio spectrum exhibits no such break, we will have to search for an alternative reason for the lack of cooling break (as we discussed above, a low value for $\\gamma_{min}$, manifested through a lack of low frequency break, requires a higher value for $B_{eq}$, which in turns produces a cooling break below the optical, as can be seen from the thin solid line in  Figure \\ref{fig:sed}). Distributed reacceleration is a very plausible candidate and it has been claimed to explain the optical emission of 3C 445 (e.g. Prieto et al. 2002).  \\subsection{What is the X-ray emission mechanism? \\label{x}}  Any interpretation of the X-rays must take into account ($i$) the level and spectrum of the X-ray emission,  ($ii$) the apparent upstream `nesting'  of the X-ray emission into the hollow part of the east-west arc of radio--optical emission, and  ($iii$) the mostly westward displacement of the X-ray emission relative to the peak of the radio - NIR 1  hotspot. Also, because of  the moderate hotspot to counterhotspot flux ratio in the radio and optical, it should not invoke significant beaming for the radio-optical emitting plasma. These considerations automatically exclude the possibilities of one zone EC/CMB \\citep{tavecchio05} and SSC in equipartition, because in these models both the SSC and EC/CMB emission is cospatial with the radio--optical hotspot emission and (see Fig. \\ref{fig:sed}) is much lower than the X-ray detected flux).  A displacement between the X-ray and radio emission is predicted in the case of UC emission from a decelerating flow, in which freshly accelerated relativistic electrons from the fast base of the flow upscatter to high energies the radio photons produced in  the downstream slow part of the flow by electrons that have cooled radiatively \\citep{georganopoulos03}. This version of the UC emission from a decelerating flow is  not favored, however, as it predicts that the optical and the X-rays are co-spatial and that the radio emission is shifted downstream. This is because optical emission can only be produced by the fast, upstream part of the flow, where the freshly accelerated electrons are found, the same place from which the UC X-rays are produced. Also, the model predicts the optical to be more beamed than the radio, something not supported by the very similar hotspot to counter-hotspot flux ratio in radio and optical. Finally, the model uses as seed photons the synchrotron photons produced in the source, making the implicit assumption that these photons dominate  the local photon energy density. This is not the case, as can be seen in Figure \\ref{fig:seeds}.    The fact that the X-ray emission is displaced from the radio-optical hotspot requires the X-ray emitting electrons to also be displaced from the radio-optical emitting electrons. If the X-ray emission is of IC nature, these electrons will experience the seed photon field of Figure \\ref{fig:seeds}, provided they are located at a distance from the radio-optical hostspot not much greater than the radio-optical hotspot size.  This immediately  excludes the IR-optical photons as seed photons for the X-ray emission, because this would require the presence of a powerful component of IC emission due to CMB photons upscattered in $\\sim$ the optical band,  co-located with the X-ray component. This is not observed. We therefore reach the conclusion that {\\sl if the X-ray emission is due to a particle  population located upstream of those in other bands, then the seed photons that these electrons upscatter must be the CMB. }  \\begin{figure} \\epsscale{1.1} \\plotone{3C445seeds.eps}  \\caption{The energy density at the location of the southern hotspot. The straight line represents the photon energy density due to the radio optical-emission of the hotspot, while the blackbody is the CMB.}  \\label{fig:seeds} \\end{figure}  A model that produces co-spatial radio-optical synchrotron emission and EC/CMB X-ray emission shifted upstream, has been proposed by \\cite{georganopoulos04} to address such  displacements observed in the large scale jets of quasars. The model assumes that distributed particle acceleration offsets radiative losses. In this model,  a relativistic decelerating flow results in an increase of the magnetic field and electron density at the slow downstream part of the flow, increasing  the synchrotron emissivity. At the same time, the faster upstream part of the flow experiences a higher CMB comoving energy density $U_{CMB}\\propto \\Gamma^2$, where  $\\Gamma$ is the bulk Lorentz factor of the flow (see Figure 3 of Georganopoulos \\& Kazanas 2004).  In this scenario the only displacement observed is along the flow axis. In our case, however, besides the general upstream shift of the X-rays relative to the radio--NIR emission, we note that the peak of the radio--NIR 1 emission is shifted  by $\\sim 2\" $ to the east relative to the center of the X-rays ($1\"$ corresponds to $1.07$ Kpc). Therefore, for this model to still be viable, the flow needs to bend to the east after producing the X-ray emission.   To demonstrate how this  could reproduce the observed upstream X-ray emission, we plot in Figure \\ref{fig:sedX}  the emission that would result from plasma in equipartition moving with a bulk Lorentz factor $\\Gamma=4$ at an angle to the line of sight $\\theta=14^\\circ$ and carrying the same power as that injected in the radio-optical hotspot ($L_{kin}=1.3 \\times 10^{44} $ erg s$^{-1}$). Although such a small angle to the line of sight, bordering angles typical of blazars is admittedly uncomfortable, it cannot be excluded for this broad line FR II radio galaxy. To reproduce the shift of the radio-optical emission in NIR 1 relative to location of the X-ray emission, the flow must bend to the east, forming an angle of $30^\\circ$ to the line of sight. The projected physical distance between the X-ray and radio-optical components is 4.6 kpc. Bends in the flow and velocity gradients at the hotspots, such as the one suggested here, are  seen in numerical simulations of relativistic flows (e.g.  Aloy et al. 1999).  Deceleration from $\\Gamma=4$ to the subrelativistic speed of the radio-optical hotspot could be achieved by a series of oblique shocks that could also aid electron re-acceleration. This latter possibility could be tested by future, high-resolution radio imaging of the hotspot. If indeed this is the X-ray emission mechanism, we expect that due to relativistic dimming, the X-ray emission from the counter hotspot would be undetectable, even for very deep {\\sl Chandra} exposures. On the assumption of hotspot-counter hotspot symmetry, X-ray detection of the counter hotspot would exclude this last alternative for an inverse Compton interpretation of the X-rays.   \\begin{figure} \\epsscale{1.10} \\plotone{3C445X.eps} \\caption{  X-rays due to EC/CMB (long dash line). The emission is assumed to come from plasma moving relativistically (bulk Lorentz factor $\\Gamma=4$) at an angle $\\theta=14^\\circ$ to the line of sight. To reproduce the observed displacements we assume that the flow decelerates and  bends to $\\sim \\theta=30^\\circ$ and terminates $\\sim 4.6$ Kpc downstream, producing the radio optical emission (solid line). The radio-optical emission coming from the X-ray spot is much weaker (short dash line). The data points are the same plotted in Figure \\ref{fig:sed}.} \\label{fig:sedX} \\end{figure}    An alternative model, which avoids the uncomfortably small jet angle to the line of sight and the relatively high Lorentz factor of the X-ray emitting plasma needed for the above interpretation, together with the arguments against strong beaming in other sources (see \\S \\ref{intro}), is to generate the X-ray emission via synchrotron radiation from a high energy population ($\\gamma\\sim 10^8$) of electrons accelerated upstream of the radio-optical hotspot.  A possible way to produce this electron population upstream the radio-optical emitting region is through acceleration in the reverse shock of a reverse-forward shock structure (in this picture the radio-optical comes from electron acceleration in the forward shock; Kataoka et al. (2008) proposed it, motivated by the upstream, relative to the radio, X-ray emission seen in both hotspots and jets of 3C  353).  This scheme, however, as Kataoka et al. (2008) note,  does not address the reason the two shocks accelerate electrons at different energies, with the reverse shock consistently reaching higher electron energies.   If the X-ray emission of the southern hotspot is synchrotron, then the observed emission in the {\\it Chandra} band must lie close to the peak of its $\\nu f_\\nu$ emission. This is because the X-ray photon index is $\\sim 2$ and the luminosity of this component at optical energies must  be below the level of the optical emission detected downstream. The emission could in principle extend to energies higher than the {\\it Chandra} band but this would require the presence of unrealistically energetic electrons.) Because cooling of the X-ray emitting electrons is severe, even if we only consider the CMB photon energy density, these electrons are suffering strong cooling, which must be  balanced by continuous reacceleration.  Finally, an account of most of the phenomenology may rest with the possibility that the entire radio-NIR-X-ray emission is synchrotron, and that the jet beam moves laterally with time to the west, implying that the maximum X-ray emission is the most recent and naturally further displaced from the AGN core. For this same reason (age) there is less IR and almost no radio emission in along this direction (the corresponding electron radiative times are longer). The fact that the radio emitting electrons have the longest radiative lifetime would then explain the displacement of the maximum emission at this frequency (and of the IR) to the East  and the absence of X-rays in the same region (the X-ray emitting electrons have all cooled to lower energies). There is some indication of such a motion from the fact that the AGN core and the two lobes do not all line on a straight line.  To conclude, both the EC/CMB from a decelerating flow  and synchrotron interpretation of the X-rays face important problems, although the EC/CMB model is more constrained and, therefore, easier to falsify, while the synchrotron interpretation  can reproduce any observed emission by introducing additional electron populations as needed. A purely spectral discrimination between the synchrotron and inverse-Compton models is not possible, as both can be made compatible with various X-ray slopes as well as the forms for the 'valley' in between the X-ray and radio-IR components (see e.g., Uchiyama et al. 2006, 2007; Jester et al. 2007; and  Hardcastle et al. 2004).  The synchrotron interpretation, however, does not require beaming and is the only alternative among those examined that could produce detectable X-ray emission  from the counter-hotspot. Therefore, detection of bright X-ray emission from the counter hotspot in future X-ray observations would rule out any reasonable IC models for the hotspots of 3C 445. An additional test of the nature of the X-ray emission could come from future X-ray polarimeters like GEMS: while  the synchrotron emission is expected to be highly polarized, the EC/CMB should produce negligible polarization (Begelman \\& Sikora 1987, McNamara et al. 2009, Uchiyama \\& Coppi in prep.). An approach that can yield results with our current observational capabilities is based on HST UV polarimetry: if the far UV emission is shown to be the low energy tail of the X-ray component as is the case with 3C 273 (e.g. Jester et al. 2007), then UV HST imaging polarimetry of the hotspots  will distinguish between the EC/CMB and synchrotron mechanisms (e.g., McNamara et al. 2009, Uchiyama \\& Coppi, in prep.)."
    },
    "0910/0910.1815_arXiv.txt": {
        "abstract": "{} {We investigate the physical and chemical conditions in a typical star forming region, including   an unbiased search for new molecules in a spectral region previously unobserved.} { Due to its proximity, the Orion KL region offers a unique laboratory of  molecular astrophysics in a chemically rich, massive star forming region. Several ground-based spectral line surveys have been made, but due to the absorption by water and oxygen, the terrestrial atmosphere is completely opaque at frequencies around 487 and 557 \\mbox{GHz}. To cover these frequencies we   used the Odin satellite to perform a spectral line survey in the frequency ranges 486\\,--\\,492 GHz and 541\\,--\\,577 \\mbox{GHz}, filling the gaps between previous spectral scans. Odin's high main beam efficiency, $\\eta_{\\mathrm{mb}}$\\,=\\,0.9, and observations performed outside the atmosphere make our intensity scale very well determined. } {We observed  280 spectral lines from 38 molecules including isotopologues,  and, in addition, 64 unidentified lines. A few U-lines have interesting frequency coincidences such as ND and the anion SH$^-$. The beam-averaged emission is dominated by CO, H$_2$O, SO$_2$, SO, $^{13}$CO and CH$_3$OH. Species with the largest number of lines are CH$_3$OH, (CH$_3)_2$O, SO$_2$, $^{13}$CH$_3$OH, CH$_3$CN and NO. Six water lines are detected including the ground state rotational transition 1$_{1,0}$\\,--\\,1$_{0,1}$  of $o$-H$_2$O, its isotopologues $o$-H$_2^{18}$O and $o$-H$_2^{17}$O,   the Hot Core tracing $p$-H$_2$O transition 6$_{2,4}$\\,--\\,7$_{1,7}$, and the 2$_{0, 2}$\\,--\\,1$_{1,1}$ transition of HDO. Other lines of special interest are the 1$_{0}$\\,--\\,0$_{\\,0}$  transition of NH$_3$ and its isotopologue $^{15}$NH$_3$. Isotopologue abundance  ratios of D/H, $^{12}$C/$^{13}$C, $^{32}$S/$^{34}$S, $^{34}$S/$^{33}$S, and $^{18}$O/$^{17}$O are estimated. The  temperatures, column densities and abundances in the various subregions are estimated, and we find very high gas-phase abundances of H$_2$O, NH$_3$, SO$_2$, SO, NO, and CH$_3$OH. A comparison with the ice inventory of ISO sheds new light on the origin of the abundant gas-phase molecules. } {} ",
        "introduction": " ",
        "conclusions": "We have in our column density and molecular abundance calculations in all cases treated the various Orion KL subregions, probed by the large Odin antenna beam, as homogeneous sources having specified average temperatures, densities and equivalent circular sizes and beam-filling factors.  However,   the \\emph{High Velocity Outflow} is known to be bi-polar with a FWHP size of 60\\,--\\,70\\arcsec, as estimated from simultaneous Odin mapping of the H$_2$O and CO \\mbox{$J$\\,=\\,5\\,--\\,4} brightness distributions (Olofsson  \\etal~\\cite{Olofsson03},  Hjalmarson  \\etal~\\cite{Hjalmarson05}). As seen from the radiative transfer equation, the H$_2$O  excitation temperature corresponding to a size of 70\\arcsec~is only 26 K, while a temperature of 72 K as found by Wright \\etal~(\\cite{Wright 2000})  corresponds to a size of  only 32\\arcsec. This indicates a  very clumpy  H$_2$O  brightness distribution,  and that this source is filled with approximately one fourth of water emission. Similar results are obtained investigating the size-temperature relation of CO. In fact, the appearance of the HVF may be similar to the clumpy, finger-like emission seen in shock-excited H$_2$ (Salas \\etal~\\cite{Salas})\\footnote{See also www: http://subarutelescope.org/Science/press-release/99$\\emptyset$1/OrionKL\\_3$\\emptyset\\emptyset$.jpg.}.  The \\emph{Low Velocity Outflow} has a NE-SW elongated structure, roughly orthogonal to the HVF, which also must have small scale structure (Genzel \\etal~\\cite{Genzel81}; Greenhill \\etal~\\cite{Greenhill98}). If we guide ourselves by the optically thick HDO lines observed by Pardo \\etal~(\\cite{Pardo}) and use a HPBW size of 15\\arcsec~for H$_2^{18}$O, the corresponding excitation temperature is only 40 K. The size corresponding to 72 K is 10$\\arcsec$ which suggests that this source is filled with about one half of radiating gas-phase water. Similar results are found for all optically thick outflow species.  The \\emph{Compact Ridge}  and \\emph{Hot Core} size-temperature relations are more consistent with the assumed values, although we know from Beuther \\etal~(\\cite{Beuther}) that both sources are very clumpy. This might be caused by their much smaller size as compared to the outflows, where the clumping affects larger scales. As seen from the CH$_3$OH Fig.   \\ref{SourceSizeCH3OH} (on-line material) and Fig. \\ref{SourceSize 2comp}, the source size of the CR also varies with the upper state energies of the lines. This might indicate considerable temperature and density variations within the CR.  Considering the uncertainties discussed above and in \\mbox{Sect. \\ref{section Theory}} and \\ref{uncertainties appendix} (Appendix), the striking agreement with B87 and S95 strengthens our confidence of our results."
    },
    "0910/0910.1738_arXiv.txt": {
        "abstract": "In this letter we present the result of the cross correlation between the 4th INTEGRAL/IBIS soft gamma-ray catalog, in the range 20-100 keV, and the Fermi LAT bright source list  of objects emitting in the 100 MeV - 100 GeV range. The main result is that only a minuscule part of the more than 720 sources detected by INTEGRAL and the population of 205 Fermi LAT sources are detected in both spectral regimes. This is in spite of the mCrab INTEGRAL sensitivity for both galactic and extragalactic sources and the breakthrough, in terms of sensitivity, achieved by Fermi at MeV-GeV energies. The majority of the 14 Fermi LAT sources clearly detected in the 4th INTEGRAL/IBIS catalog are optically identified AGNs (10) complemented by 2 isolated pulsars (Crab and Vela) and 2 High Mass X-Ray Binaries (HMXB, LS I +61$^{\\circ}$303 and LS 5039). Two more possible associations have been found: one is 0FGL J1045.6-5937, possibly the counterpart at high energy of the massive colliding wind binary system Eta Carinae, discovered to be a soft gamma ray emitter by recent INTEGRAL observations and 0FGL J1746.0-2900 coincident with IGR J17459-2902, but still not identified with any known object at lower energy. For the remaining 189 Fermi LAT sources no INTEGRAL counterpart was found and we report the 2$\\sigma$ upper limit in the energy band 20-40 keV. ",
        "introduction": "In the last few years, our knowledge of the soft gamma-ray sky (E$>$20 keV) has greatly improved thanks to the observations performed by the imager IBIS (Ubertini et al. 2003) onboard INTEGRAL (Winkler et al. 2003) and  by the BAT telescope (Barthelmy et al. 2005) onboard Swift (Gehrels et al. 2004). In particular, IBIS is surveying the 20--100 keV sky with a sensitivity better than a mCrab and a point source location accuracy of the order of a few arcmin. The 4th IBIS catalog has been just released, listing 723 soft gamma-ray sources (Bird et al. 2009). It represents an extension both in exposure and sky coverage with respect to the previous catalog (Bird et al. 2007). The IBIS soft gamma-ray  sky is surprisingly dynamic and diverse, with many different types of objects such as Active galactic Nuclei (AGNs), galactic X-ray binaries, Cataclysmic Variables (CVs.) and also newly  discovered classes of sources such as  Supergiant Fast X-ray Transients (SFXTs, Sguera et al. 2005,2006) and highly absorbed supergiant high mass X-ray binaries (Walter et al. 2006).  Recently, INTEGRAL and Swift have been joined in operation by the Fermi gamma-ray satellite launched in June 2008, providing a coverage of the sky over 10 orders of magnitude in energy, from keV to GeV. Onboard this satellite, the Large Area Telescope (LAT, Atwood et al. 2009) is an imaging high energy gamma-ray telescope (20 MeV -- 300 GeV) which provides, with respect to its predecessor EGRET, superior angular resolution, sensitivity, field of view and observing efficiency. These improved capabilities have allowed a catalog of  205 highly significant detections ($>$10$\\sigma$) after three months of an all sky survey (Abdo et al. 2009a). Many Fermi LAT sources  have been identified as AGNs (121) and  pulsars (30), but interestingly there are also 2 High Mass X-ray Binaries (HMXBs) while the remaining objects have no obvious counterparts. To identify those sources IBIS is particularly suited since it provides a hard X/soft gamma-ray point source location accuracy of the order of an arcmin as well as information close to the MeV band. ",
        "conclusions": "From our analysis it appears that MeV/GeV Fermi sources are not commonly associated with  IBIS sources in the range 20 -- 100 keV. The handful of objects common to both surveys comprise so far mainly FSRQ and BL Lac, with no XRB apart from the two microquasars. Equally absent are the AXP which are strong emitters in the  keV-MeV range with a total energy rising in $\\nu$F$\\nu$ and no cut-off detected up to a few hundreds of keV (Kuiper and Hermsen 2009), implying some kind of switch-off mechanism in the MeV regime. Similarly, SGR and Magnetars, detected even in quiescence mode by IBIS (Rea et al, 2009) and among the brightest sources when flaring (Kouveliotou et al. 1999, Gotz et al. 2006, Israel et al. 2008) are not detected so far by Fermi/LAT.   \\begin{landscape} \\begin{table*}[t!] \\begin{center} \\caption {Spatial correlation between Fermi LAT and IBIS sources.\\tablenotemark{1}} \\begin{tabular}{cccccccc} \\hline \\hline Fermi        & RA      & Dec     & IBIS           & IBIS         & distance   &                                                               & ID source \\\\ (0FGL)       & (deg)   &(deg)    & RA (deg)             & Dec  (deg)         &   (arcmin) & $\\frac{distance}{\\sqrt{Fermi R_{1\\sigma}^2 + Ibis R_{1\\sigma}^2}}$  &          \\\\ \\hline J1229.1+0202 & 187.287 & 2.045   & 187.279        & 2.049        & 0.533      & 0.258        & 3C 273\\\\ J1325.4-4303 & 201.353 & -43.062 & 201.363        & -43.021      & 2.5        & 0.332       &  Cen A \\\\ J1653.9+3946 & 253.492 & 39.767  & 253.488        & 39.753       & 0.85       & 0.410     & Mkn 501\\\\ J2202.4+4217 & 330.662 & 42.299  & 330.677        & 42.293       & 2.01       & 0.481  & BL Lac\\\\ J2254.0+1609 & 343.502 & 16.151  & 343.489        & 16.149       & 0.75       & 0.558    & 3C 454.3\\\\ J0240.3+6113 & 40.093  & 61.225  &  40.119     & 61.240       & 1.16       & 0.562  & LS I+616 303\\\\ J0320.0+4131 & 50.000  & 41.524  & 49.967      & 41.532       & 1.55       & 0.645   &  NGC 1275 \\\\ J1256.1-0547 & 194.034 & -5.8   & 194.044     & -5.770       & 1.88       & 0.827 &  3C 279 \\\\ J1104.5+3811 & 166.137 & 38.187 & 166.119     & 38.207       & 1.46       & 1.067   & Mkn 421\\\\ J1833.4-2106 & 278.37  & 21.103 & 278.415     & -21.063      & 3.46       & 1.180  & PKS 1830-211\\\\ J0534.6+2201 & 83.653  & 22.022 & 83.629      & 22.017       & 1.36       & 1.190    & Crab \\\\ J1826.3-1451 & 276.595 & -14.86 & 276.525     & -14.847      & 4.13       & 1.406   &  LS 5039\\\\ J0835.4-4510 & 128.865 & -45.17 & 128.831 & -45.179       & 1.53       & 1.437  & Vela Pulsar\\\\ J1746.0-2900 & 266.506 &-29.005 & 266.485 &-29.043           & 2.51       & 1.472   & IGR J17459-2902\\\\ J0036.7+5951 & 9.177   & 59.854 & 8.964   & 59.83         & 6.56       & 1.802 & 1ES 0033+59.5 \\\\ J1045.6-5937 & 161.409 & -59.631& 161.206 & -59.704       & 7.5        & 2.073 &  Eta Carinae \\\\ \\hline \\hline \\end{tabular} \\end{center} \\tablenotetext{1}{Note: the sixth column is the distance between the LAT and IBIS source positions while the seventh column is this distance divided the square root of the quadratic sum of their corresponding error radii.}  \\end{table*} \\end{landscape}   \\begin{figure}[t!] \\plotone{Correl1a.eps} \\caption{Top: the number of IBIS sources associated with Fermi (solid line) and fake (dashed) objects. Bottom: the difference between the two upper curves revealing a correlation for a small number of sources:} \\end{figure}   \\begin{figure}[t!] \\plotone{eta_18_60.eps} \\caption{IBIS 18-60 KeV significance map ($\\sim$ 2.3 Ms exposure time) superimposed on the Fermi LAT errror circle of the unidentified source 0FGL J1045.6-5937. It is evident a clear IBIS detection at $\\sim$ 5.5$\\sigma$. The optical position of Eta Carinae is marked by a black cross point.} \\end{figure}   \\begin{figure}[t!] \\plotone{fig4.eps} \\caption{Gamma-ray flux  (100 MeV - 1 GeV) of each Fermi  source as a function of the corresponding 20--40 keV IBIS flux. The coloured points refer to the IBIS detections, specifically red points are blazars, dark blu are pulsars, green are HMXBs, yellow is Eta Carinae and finally light blu is IGR J17459-2902. The black points refer to IBIS non detection (2$\\sigma$ upper limit). } \\end{figure}   \\begin{landscape} \\begin{table*}% \\begin{center} \\caption {2$\\sigma$ upper limit (20-40 keV) for the 189 Fermi LAT sources with no IBIS counterpart\\tablenotemark{2}} \\begin{tabular}{cccccccc} \\hline \\hline Fermi         &  upper limit    &     Fermi   &     upper limit   & Fermi  &  upper limit &    Fermi  &   upper limit     \\\\ (0FGL)       &  20-40 KeV       &   (0FGL)   &   20-40 keV        & (0FGL)&   20-40 keV      &  (0FGL)   &  20-40 keV     \\\\ &  (mCrab)       &           &   (mCrab)        &          &   (mCrab)      &      &   (mCrab)       \\\\ \\hline J0007.4+7303      &    0.3084       &     J0654.3+5042   &       2.3834       &     J1331.7-0506   &      0.5776      &  J1834.4-0841     &    0.2302   \\\\ J0017.4-0503      &    2.3976       &    J0700.0-6611     &      0.5454          &      J1333.3+5058  &   0.6076     &      J1836.1-0727 &     0.2332    \\\\ \\hline \\hline \\end{tabular} \\end{center} \\tablenotetext{2}{Note. Table 2 is published in its entirety in the electronic edition of the Astrophysical Journal Letters. A portion is shown here for guidance regarding its form and content. The first (and odd) column(s) list the Fermi LAT bright sources emitting in the 100 MeV - 100 GeV range (Abdo et al., 2009); the second (and even) column(s) report the Integral/IBIS 2sigma upper limit in the energy band 20-40 keV. The IBIS fluxes are in milliCrab = 7.6 x 10$^{-12}$erg cm$^{2}$s in the 20-40 KeV range.} \\end{table*} \\end{landscape}"
    },
    "0910/0910.1801_arXiv.txt": {
        "abstract": "{} {With the goal of providing constraints on the nature of the progenitors of core-collapse (CC) supernovae (SNe), we compare their radial distribution within their spiral host galaxies with the distributions of stars and ionized gas in the spiral disks.} {SNe positions are taken from the Asiago catalog for a well-defined sample of 224 SNe within 204 host galaxies. The SN radial distances are estimated from the deprojected separations from the host galaxy nuclei, and normalized both to the $25^{\\rm th} {\\rm mag~arcsec}^{-2}$ blue-band isophotal radius and (for the first time) to the statistically-estimated disk scale length.} {The normalized radial distribution of all CCSNe is consistent with an exponential law, as previously found, with a possible depletion of CCSNe within one-fifth of the isophotal radius (not seen with scale-length normalization). There are no signs of truncation of the exponential distribution of CCSNe out to 7 disk scale lengths. The scale length of the  distribution of type~II SNe appears to be significantly larger than that of the stellar disks of their host galaxies, but consistent with the scale lengths of Freeman disks. SNe~Ib/c have a significantly smaller scale length than SNe~II, with little difference between types Ib and Ic. The radial distribution of type Ib/c SNe is more centrally concentrated than that of the stars in a Freeman disk, but is similar to the stellar disk distribution that we infer for the host galaxies. All CCSN subsamples are consistent with the still uncertain distribution of \\ion{H}{ii} regions. The scale length of the CCSN radial distribution shows no significant correlation with the host galaxy morphological type, or the presence of bars. However, low luminosity as well as inclined hosts have a less concentrated distribution (with the statistical scale-length normalized radial distances) of CCSNe, which are probably a consequence of metallicity and selection effects, respectively.} {The exponential distribution of CCSNe shows a scale length consistent with that of the ionized gas confirming the generally accepted hypothesis that the progenitors of these SNe are young massive stars. Given the lack of correlation of the normalized radial distances of CCSNe with the morphological type of the host galaxy, we conclude that the more concentrated distribution of SNe~Ib/c relative to SNe~II must arise from the higher metallicity of their progenitors or possibly from a shallower initial mass function in the inner regions of spirals.} ",
        "introduction": "Most SNe can be assigned to two main physical classes \\citep[e.g.,][]{Turatto03,Tur+07}: the gravitational collapse of young massive stellar cores (Types Ib, Ic and II SNe) and the thermonuclear explosions of a white dwarf in close binary systems (Type Ia SNe). While SNe with no or weak hydrogen lines were classified into type I, it is now understood that there are actually three spectroscopically and photometrically distinct subclasses of SNe~I. Type Ia SNe are characterized by spectra with no hydrogen lines and strong \\ion{Si}{II} lines \\citep[e.g.,][]{Leibundgut2000,Liv01}. These SNe appear in galaxies of all morphological types \\citep[e.g.,][]{BBCT99,vdBLF05}. Type Ib SNe are characterized by spectra with no evident hydrogen lines, weak or absent \\ion{Si}{II} and strong \\ion{He}{I} lines. The third subclass, type Ic SNe, discovered later, shows weak or absent hydrogen, helium lines, and \\ion{Si}{II}. Type II SNe are characterized by the obvious presence of hydrogen lines. This SN type displays a wide variety of properties: type II Plateau SNe (SNe~IIP) with flat light curves in the first few months; type II Linear SNe (SNe~IIL) with a rapid, steady decline in the same period; the narrow-lined SNe (SNe~IIn), dominated by emission lines with narrow components, a sign of energetic interaction between the SN ejecta and the circumstellar material; and type IIb SNe (SNe~IIb), a transitional type, which have early time spectra similar to SNe~II and late time spectra similar to SNe~Ib/c \\citep[e.g.,][]{Filippenko+90,Filippenko+94}.  With the exception of SNe~Ia, all types of SNe are rare in early-type galaxies \\citep[e.g.,][]{vdBLF05}. Surprisingly, among morphologically classified hosts of SNe~Ib/c and SNe~II \\citet{HPM+2008} found 22 cases where the host has been classified as an Elliptical or S0 galaxy. However, all the early-type hosts of SNe~Ib/c and SNe~II display independent indicators of recent star formation due to merging or gravitational interaction.  According to theoretical models, the progenitors of SNe~Ib/c are massive O type stars that have lost most or all of their hydrogen (and perhaps their helium) envelopes, either by strong winds as in Wolf-Rayet stars \\citep[e.g.,][]{Eldridge+04}, or through the transfer of material to a binary companion via Roche lobe overflow \\citep[e.g.,][]{Heger+03}. Progenitors of SNe~II are massive stars that retain their hydrogen envelopes \\citep[e.g.,][]{WW86,Hamuy03rev}.  Many observational studies confirm these theoretical interpretations of these classes of core-collapse SNe (CCSNe). The rates of all CCSN types depend on the morphology of the host galaxies. The rate of SNe~Ib/c and II per unit stellar mass increases by factors of 3 and 5 respectively from early- to late-type spirals host galaxies \\citep[e.g.,][]{Mannucci+05}. Given the short lifetime of their massive progenitors, the rate of CCSNe in host galaxies directly traces the current star formation rate \\citep[e.g.,][]{Mal03}. Conversely, when the star formation rate is known, it can be used to verify the consistency of the progenitor scenario \\citep[e.g.,][]{HB06}.  The observation of progenitor stars of CCSNe in archival pre-explosion images provides a direct test of the theoretical predictions. Since massive evolved stars are the most luminous objects in a galaxy, the progenitors of CCSNe should be directly detectable on pre-explosion images of nearby host galaxies. The advent of data archives of large telescopes with high image quality, most importantly that of the Hubble Space Telescope, has allowed to extend the progenitor detection to distances larger than 20 Mpc \\citep[e.g.,][]{vD+03,Mau+05}. The observations have generally confirmed  the theoretical predictions mentioned above \\citep[e.g.,][]{Crockett+07,Pastorello+07,Smartt+08}.  The spatial distribution of SNe in host galaxies provides another strong constraint on the nature of SN progenitors. Various studies show that CCSNe are tightly connected to the disk \\citep[e.g.,][]{JM63} and to the spiral arm structure (\\citeauthor{JM63}; \\citealp{MvdB76,BTF94,vDHF96}), which dramatically differs from the SN~Ia distribution \\citep[e.g.,][]{FoSch08}. In addition, CCSNe are well associated with star-forming sites: OB-associations, and \\ion{H}{ii} regions \\citep{BTF94,vDHF96,TBP01}. SNe~Ib/c show a higher degree of association with \\ion{H}{ii} regions than those of type II SNe, suggesting that they arise from a higher mass progenitor population than the SNe~II \\citep{AndeJa08}. Most recently, \\citet{AndeJa09} found that the radial positions of a sample of 177 CCSNe closely follow the radial distribution of H$\\alpha$ emission, implying that these SNe are excellent tracers of star formation within galaxies.  One can better constrain the nature of the progenitors of CCSNe by considering quantitative measures of their radial distribution within their host galaxies, after correcting for the inclination of the disk. In their pioneering study, \\citeauthor{JM63} found a rapidly decreasing surface density distribution of SNe, except for an important lack of SNe in the central regions of spiral galaxies. The surface density of SNe is known to be exponential \\citep{BCC75,VZ77}, although two studies \\citep{IK75,GKK80} suggested a ring like distribution. \\citeauthor{BCC75} quantified the exponential to have a slope that amounts to a scale length of $0.46\\pm0.03$ times the optical radius, while \\citeauthor{VZ77} found a scale length of $0.12\\pm0.01$ times the optical radius. The difference in scale length arises from the different definitions of optical radius. All SN types, even the SNe Ia, were included in these early studies. Using a large sample of CCSNe (of known types), i.e. 74 SNe~Ib/c and SNe~II, \\citet{vdB97} found  a scale length of $0.22\\,R_{25}$ (with no error bar given), where  $R_{25}$ is the isophotal radius for the blue-band surface brightness of $\\mu_B=25\\,\\rm mag\\,arcsec^{-2}$. Also, \\citet{Bar+92} found that their 99 SNe~II have an exponential distribution with scale length that amounts to $(0.27\\pm0.08)\\,R_{25}$, extending far beyond the optical radius of galaxies. Moreover, \\citeauthor{Bar+92} noticed a sharp decrease in slope of both SNe~II and SNe~Ib at $1.4\\,R_{25}$.  It has rapidly become apparent that SNe~Ib/c are more centrally concentrated within galaxy disks than are SNe~II. \\citeauthor{Bar+92} noticed that their 22 SNe~Ib show a non-exponential distribution that increases slope with radius so that the scale length varies from $(0.15\\pm0.02)\\,R_{25}$ inside $0.5\\,R_{25}$ to $(0.10\\pm0.02)\\,R_{25}$ between 0.5 and $1.0\\,R_{25}$. The more centrally concentrated distribution of SNe~Ib/c in comparison to that of SNe~II was also noticed by \\citet{vdB97} and \\citet{Wang+97}, but, in contrast to the study of \\citeauthor{Bar+92}, the difference with the distribution of SNe~II was not statistically significant.  Finally, \\citet{TPB04} found that, while in the central regions, SNe~Ib/c are more concentrated than are the SNe~II, in the outer regions the two distributions are similar. This result is at odds with the results of \\citet{Bar+92}, who found increasingly dissimilar slopes for the surface densities of SNe~II and of SNe~Ib/c in the outer regions of spirals.  Additional insight is obtained by comparing the radial distribution of CCSNe in galaxies of different activity or environment. \\citet{PT90} found that their sample of 8 SNe~II and SNe~Ib within galaxies hosting AGN were significantly more radially concentrated in their galaxy hosts than analogous CCSNe in galaxies without active nuclei. \\citet{Petr+05} studying a sample of 12 SNe~II and SNe~Ib/c in galaxies hosting AGN,  confirmed the above result and found that SNe~Ib/c in active/star-forming galaxies are more centrally concentrated than are the SNe~II, but given the small sample, this difference was not statistically significant. These results were confirmed with larger samples of CCSNe by \\cite{Hak08}, who used both one-dimensional and multivariate statistics.  The locations of SN explosions in multiple systems have also been studied. In interacting galaxies, CCSNe are not preferentially located towards the companion galaxy \\citep{Petr+95,Nav+01}. Similarly, the azimuthal distributions inside the host members of galaxy groups are consistent with being isotropic (\\citealp{Nav+01}).  In this paper, we make use of a considerably larger sample of CCSNe to reanalyze the distribution of deprojected SN radii normalized both to $R_{25}$ and, for the first time, to indirect estimates of individual stellar disk scale lengths. We make statistical comparisons between the radial distributions of CCSNe and the distribution of blue light and \\ion{H}{ii} regions in the disks of spiral galaxies.  The plan of the paper is as follows. The samples of CCSNe and host galaxies are presented in Sect.~\\ref{sample}. We discuss the normalizations in Sect.~\\ref{norm}. The results, given in Sect.~\\ref{results}, are discussed and summarized in Sect.~\\ref{discus}. Throughout this article we have assumed a value of $H_0=75 \\,\\rm km \\,s^{-1} \\,Mpc^{-1}$ for the Hubble constant. ",
        "conclusions": " \\begin{enumerate} \\item The surface density of CCSNe in all studied subsamples (different CCSN type or galaxy host type or luminosity or inclination) falls exponentially with relative radius $R_{\\rm SN} / R_{25}$ and  $R_{\\rm SN} /h$, with no signs of truncation out to 7 disk scale lengths; \\item The radial distribution of CCSNe is not significantly affected by the host galaxy type, or by the presence of a bar in the host; \\item The radial distribution of CCSNe, measured with scale-length normalization, is more concentrated in high luminosity host galaxies; \\item The radial distribution of SNe~Ib/c is significantly more concentrated than the Freeman disk distribution and consistent with the radial distribution of \\ion{H}{ii} regions; \\item The radial distribution of SNe~II is consistent with both the Freeman disk and the \\ion{H}{ii} regions distributions, but significantly less concentrated than the host disks; \\item There is a small lack of CCSNe within one-fifth of the isophotal radius ($R_{\\rm SN} < 0.2\\,R_{25}$), not well visible with scale-length normalization; \\item The radial distribution of type Ib/c SNe in their host galaxies is more centrally concentrated than that of type II SNe, the ratio of scale lengths is $0.77\\pm0.03$, probably because of a metallicity effect. \\end{enumerate}   It would be worthwhile to extend these analyses, by comparing the radial distribution of CCSNe with recent accurate measures of the distribution of molecular gas (i.e. CO), as well as ionized gas (see compilation by \\citealp{BPBG03}) and even neutral gas, and by extending the analysis to two dimensions (extending previous analyses such as those by \\citealp{JM63}, \\citealp{BTF94}, \\citealp{vDHF96}, and \\citealp{TBP01}, making use of the much larger CCSN sample used here)."
    },
    "0910/0910.4109_arXiv.txt": {
        "abstract": "We study the quark deconfinement phase transition in hot $\\beta$-stable hadronic matter. Assuming a first order phase transition, we calculate the enthalpy per baryon of the hadron-quark phase transition. We calculate and compare the nucleation rate and the  nucleation time due to thermal  and quantum nucleation mechanisms. We compute the crossover temperature above which thermal nucleation dominates the finite temperature quantum nucleation mechanism. We next discuss the consequences for the physics of proto-neutron stars. We introduce the concept of limiting conversion temperature and critical mass  $M_{cr}$  for proto-hadronic stars, and we show that  proto-hadronic stars with a mass $M < M_{cr}$ could survive the early stages of their evolution without decaying to a quark star. ",
        "introduction": "In the last few years there has been a growing interest in the study of the nucleation process of quark matter (QM) in the core of massive neutron stars. In particular, it has been shown \\cite{be03,bo04,drago04,lug05,blv07,bppv08,dps-b08,bambi08} that above a threshold value of the central pressure  a pure hadronic compact star (HS) is metastable to the decay (conversion) to a quark star (QS) ({\\it i.e.} to a hybrid neutron star or to a strange star \\cite{bod71,witt84}, depending on the details of the equation of state (EOS) for quark matter used to model the phase transition \\cite{fj84,gle96,web05,haens_book}). This stellar conversion process liberates a huge amount of energy (a few $10^{53}$ erg) \\cite{grb} and it could be the energy source of some of the long Gamma Ray Bursts (GRBs).  The research reported in Refs \\cite{be03,bo04,drago04,lug05,blv07,bppv08,dps-b08,bambi08} has focused on the quark deconfinement phase transition in cold ($T = 0$) and neutrino-free neutron stars. In this case the formation of the first drop of QM could take place solely via a quantum nucleation process.  A neutron star at birth (proto-neutron star) is very hot (T =  10 -- 30 MeV) with neutrinos being still trapped in the stellar interior \\cite{BurLat86,prak97}. Subsequent neutrino diffusion causes deleptonization and heats the stellar matter to an approximately uniform entropy per baryon $\\tilde {S}$ =1 -- 2 (in units of the Boltzmann's constant $k_B$). Depending on the stellar composition, during this stage neutrino escape can lead the more ``massive'' stellar configurations to the formation of a black hole \\cite{bomb96,prak97}.  However, if the mass of the star is sufficiently small, the star will remain stable  and it will cool to temperatures well below 1~MeV  within a cooling time $t_{cool} \\sim$~a~few~$10^2$~s, as the neutrinos continue to carry energy away from the stellar material \\cite{BurLat86,prak97}. Thus in a proto-neutron star, the quark deconfinement phase transition will be likely triggered by a  thermal nucleation process. In fact, for sufficiently high temperatures, thermal nucleation is a much more efficient process with respect to the quantum nucleation mechanism.  Some of the earlier studies on quark matter nucleation (see {\\it e.g.}, \\cite{ho92,ho94,ol94,hei95,harko04}) have already dealt with thermal nucleation in hot and dense hadronic matter. In these studies, it was found that the prompt formation of a critical size drop of quark matter via thermal activation is possible above a temperature of about $2-3$ MeV.  As a consequence, it was inferred that pure hadronic stars are converted to quark stars within the first seconds after their birth. However, these works \\cite{ho92,ho94,ol94,hei95} reported an estimate of the thermal nucleation based on \"typical\" values for the thermodynamic properties characterizing the central part of neutron stars.   Our main objective in this and next related works, is to establish if a newborn hadronic star (proto-hadronic star) could survive the early stages of its evolution without \"decaying\" to a quark star. In the present Letter, we calculate the thermal nucleation rate of quark matter in hot ($T \\neq 0$) and neutrino-free hadronic matter using a finite temperature EOS for hadronic and quark matter. In addition, we calculate the quantum nucleation rate at finite temperature, and compare the thermal and quantum nucleation time at different temperatures and pressures characterizing the central conditions of metastable proto-hadronic compact stars. We compute the crossover temperature above which thermal nucleation dominates above the finite temperature quantum nucleation mechanism. Finally we briefly discuss some consequences for the physics of proto-neutron stars. ",
        "conclusions": "In Fig. 1 we plot the Gibbs' energies per baryon for the hadron-phase and the Q*-phase at different temperatures, T = 0, 10, 20, 30 MeV;  lines with the larger slope refer to the hadron-phase. As we see, the transition pressure $P_0$ (indicated by a full dot) decreases when the hadronic matter temperature is increased. The phase equilibrium curve  $P_0(T)$ for the hadron-quark phase transition (within the present schematic model for the EOS) is shown in Fig. 2.  As it is well known, for a first-order phase transition the derivative $dP_0/dT$ is related to the specific  ({\\it i.e.} per baryon) latent heat ${\\cal Q}$  of the phase transition by the  Clapeyron-Clausius equation \\begin{equation} \\frac{dP_0}{dT} =  - \\frac{ n_H n_{Q^*}}{n_{Q^*} - n_H } \\frac{{\\cal Q}}{T } \\label{eq:eq10} \\end {equation} \\begin{equation} {\\cal Q}  = \\tilde W_{Q^*} - \\tilde W_H = T(\\tilde {S}_{Q^*} - \\tilde {S}_H) \\label{eq:eq11} \\end {equation} where $\\tilde W_H$ ($\\tilde W_{Q^*}$)  and $\\tilde {S}_H$ ($\\tilde {S}_{Q^*}$) denote  respectively the  enthalpy  per baryon and entropy per baryon  for the hadron (quark) phase. The specific latent heat ${\\cal Q}$ and the phase numbers densities  $n_H$ and  $n_{Q^*}$ at phase equilibrium are reported in Table 1.  As expected for a first order phase transition one has a discontinuity jump in the phase number densities: in our particular case $n_{Q^*}(T,P_0) > n_H(T,P_0)$. This result, together with the positive value of ${\\cal Q}$ ({\\it i.e.} the deconfinement phase transition absorbs heat) tell us (see Eq.\\ (\\ref{eq:eq10})), that the transition temperature decreases with pressure (as in the melting of ice).  \\begin{figure}[t] \\phantom{a}\\vspace*{5mm} {\\psfig{figure=phase_diagram_B85.eps,width=7.4cm}} \\caption{Phase equilibrium curve for the hadron to quark matter phase transition.} \\label{Fig2} \\end{figure} \\begin{table} \\begin{tabular}{ccccc} \\hline $T$~~~ & ${\\cal Q}$~~~~  & $n_{Q^*}$~~~~  & $n_{H}$~~~~ & $P_0$\\\\ MeV~~~ & MeV~~~~ & fm$^{-3}$~~~~ & fm$^{-3}$~~~~ & MeV/fm$^3$ \\\\ \\hline 0~~~     &    0.00~~~~ &     0.453~~~~&  0.366~~~~  &     39.95 \\\\ 5~~~     &    0.56~~~~ &     0.451~~~~&  0.364~~~~  &     39.74 \\\\ 10~~~     &    2.40~~~~ &     0.447~~~~&  0.358~~~~  &     38.58 \\\\ 15~~~     &    5.71~~~~ &     0.439~~~~&  0.348~~~~  &     36.55 \\\\ 20~~~     &  10.60~~~~ &     0.428~~~~&  0.334~~~~  &     33.77 \\\\ 25~~~     &  17.17~~~~ &     0.414~~~~&  0.316~~~~  &     30.36 \\\\ 30~~~     &  25.44~~~~ &     0.398~~~~&  0.294~~~~  &    26.53 \\\\ \\hline \\end{tabular} \\caption{The specific  latent heat ${\\cal Q}$ and the phase numbers densities  $n_H$ and  $n_{Q^*}$ at phase equilibrium. } \\label{table} \\end{table}  In Fig. 3, we represent the energy barrier for a virtual drop of the Q*-phase as a function of the droplet radius and for different temperatures  at a fixed  pressure $P = 57$~MeV/fm$^3$.  As expected, from the results plotted in Fig. 1,  the energy barrier $U({\\cal R}, T)$ and the droplet critical radius  ${\\cal R}_c$ decrease as the matter temperature is increased. This effect favors the Q*-phase formation, and in particular increases (decreases) the quantum nucleation rate (nucleation time $\\tau_q$) with respect to the corresponding quantities  calculated at  $T=0$.  \\begin{figure}[t] \\phantom{a}\\vspace*{5mm} {\\psfig{figure=barrier_T.eps,width=7.4cm}} \\caption{Energy barrier for a virtual drop of the Q*-phase as a function of the droplet radius and for different temperatures for a pressure $P = 57$~MeV/fm$^3$.} \\label{Fig3} \\end{figure}  In Fig. 4  we plot the quantum and thermal nucleation times of the  Q*-phase  in $\\beta$-stable hadronic matter as a function of temperature and at a fixed pressure $P=57$~MeV/fm$^3$. As expected, we find a crossover temperature $T_{co}$ above which thermal nucleation is dominant with respect to the quantum nucleation mechanism.  For the case reported in Fig. 4,  we have $T_{co}=7.05$~MeV and the corresponding nucleation time is $\\log_{10}(\\tau/{\\rm s}) = 54.4$. The crossover temperature, for different values of the pressure of $\\beta$-stable hadronic matter, is reported in Tab. 2 (second column) together with the nucleation time calculated at $T = T_{co}$  (third column).  \\begin{figure}[t] \\phantom{a}\\vspace*{5mm} {\\psfig{figure=time_57.eps,width=7.4cm}} \\caption{(Color online.) Thermal ($\\tau_{th}$) and quantum ($\\tau_q$) nucleation time of quark matter (Q*-phase) in $\\beta$-stable hadronic matter as a function of temperature at fixed pressure $P= 57$~MeV/fm$^3$.  The crossover temperature is  $T_{co}=7.05$~MeV. The limiting conversion temperature for the proto-hadronic star is, in this case, $\\Theta = 10.3$~MeV, obtained from the intersection of the thermal nucleation time curve (continuous line) and the dot-dashed line representing  $\\log_{10}(\\tau/{\\rm s}) = 3$.} \\label{Fig4} \\end{figure} \\begin{table} \\begin{tabular}{ccc} \\hline $P$   & $T_{co}$ &~~$\\log_{10}(\\tau/{\\rm s}) $  \\\\ \\hline 53.98~~~~&~~  5.0~~~~~&  233.6  \\\\ 55.48~~~~&~~  6.0~~~~~&  121.3  \\\\ 56.94~~~~&~~  7.0~~~~~&    56.6  \\\\ 58.42~~~~&~~  8.0~~~~~&    16.0  \\\\ 58.85~~~~&~~  8.3~~~~~&      3.0   \\\\ \\hline \\end{tabular} \\caption{Crossover temperature $T_{co}$ (in MeV), for different fixed values of the pressure $P$ (in Mev/fm$^3$) of hadronic matter. The third column reports the logarithm of the nucleation time (in seconds) calculated at the crossover temperature. The value $8.3$ MeV defines the value of the limiting conversion temperature $\\Theta$ for a star with a central pressure $P=58.85$ MeV fm$^{-3}$.} \\label{table} \\end{table}  Having in mind the physical conditions in the interior of a proto-hadronic star \\cite{BurLat86,prak97} (see Sect. 1 of the present paper), to establish if this star will survive the early stages of its evolution without \"decaying\" to a quark star, one has to compare the quark matter nucleation time $\\tau=\\min(\\tau_q,\\tau_{th})$ with the cooling time  $t_{cool} \\sim$~a~few~$10^2$~s. If $\\tau >> t_{cool}$ then quark matter nucleation will not likely occur in the newly formed star, and this star will evolve to a cold deleptonized configuration. We thus introduce the concept of {\\it limiting conversion temperature} $\\Theta$ for the proto-hadronic star and define it as the value of the stellar central temperature $T_c$ for which the Q*-matter nucleation time is  equal to $10^3$~s. The limiting conversion temperature $\\Theta$ will clearly depend on the value of the stellar central pressure (and thus on the value of the stellar mass).  \\begin{figure}[t] \\phantom{a}\\vspace*{5mm} {\\psfig{figure=theta_B85.eps,width=7.4cm}} \\caption{(Color online.) The limiting conversion temperature $\\Theta$ for a newborn hadronic star as a function of the  central stellar pressure. Newborn hadronic stars with a central temperature and pressure located on the right side of the curve $\\Theta(P)$ will nucleate a Q*-matter drop during the early stages of their evolution, and will finally evolve to cold and deleptonized quark stars, or will collapse to  black holes. The lines labeled  $T_S$ represent the stellar matter temperature as a function of pressure at fixed entropies per baryon $\\tilde S/k_B = 1$ (dashed line) and $2$ (solid line).} \\label{Fig5} \\end{figure}  The limiting conversion temperature $\\Theta$ is plotted in Fig. 5 as a function of the stellar central pressure.  A proto-hadronic star with a central temperature  $T_c > \\Theta$ will likely nucleate a Q*-matter drop during the early stages of its evolution, and will finally evolve to a cold and deleptonized quark star, or will collapse to a black hole (depending on the particular model adopted for the matter EOS).  For an isoentropic stellar core \\cite{BurLat86,prak97}, the central temperature of the proto-hadronic star is given, for the present EOS model, by the lines labeled by $T_S$ in Fig. 5, relative to the case $\\tilde {S} = 1~k_B$ (dashed curve) and $\\tilde {S} = 2~k_B$ (continuous curve). The intersection point ($P_S,\\Theta_S$) between the two curves $\\Theta(P)$ and $T_S(P)$ thus gives the central pressure and temperature of the configuration that we denote as the {\\it critical mass} configuration of the proto-hadronic stellar sequence. The value of the gravitational critical mass $M_{cr} = M(P_S,\\Theta_S)$  and baryonic critical mass $M_{B,cr}$ are reported in Tab.~3, for three different choices of the entropy per baryon, $\\tilde {S}/k_B = 0$ (corresponding to a cold hadronic star) {\\footnote {Notice that in ref.\\cite{be03,bo04,drago04,lug05,blv07,bppv08} the critical mass for cold ($T=0$) metastable hadronic stars has been defined as the value of the gravitational mass for which the quantum nucleation time is equal to one year:  $M_{cr}(T=0) = M(\\tau_q=1~{\\rm yr})$. It is worth recalling that the nucleation time $\\tau_q$ is an extremely steep function of the hadronic star mass \\cite{be03,bo04}, therefore the exact value of $\\tau_q$ chosen in the definition of $M_{cr}(T=0)$ is not crucial (one must take a ``reasonable small'' value of $\\tau_q$, much shorter that the age of young pulsars as the Crab pulsar). We have verified that changing  $\\tau_q$ from 1~yr to $10^3$~s modifies $M_{cr}(T=0)$ by $\\sim 0.02\\%$. On the other hand, the nucleation time $\\tau=\\min(\\tau_q, \\tau_{th})$ entering in the definition of the critical mass of proto-hadronic stars $M_{cr}(\\tilde{S})$ must be comparable to the proto-hadronic star cooling time $t_{cool}$. }} , $1$ and $2$. In the same table, we also report the value of the gravitational mass ${\\cal M}$ of the cold hadronic star with baryonic mass equal to $M_{B,cr}$. This configuration is stable ($\\tau=\\infty$) with respect to Q*-matter nucleation in the case $\\tilde {S}/k_B = 2$, and it is essentially stable (having a nucleation time enormously larger than the age of the universe) in the case $\\tilde {S}/k_B = 1$. Note that these numbers are model dependent and, therefore, one must take them just as indicative values. A careful analysis is beyond the scope of the present work and it will be addressed in a future work.  \\begin{table} \\begin{tabular}{cccc} \\hline $\\tilde {S}/k_B$ & $M_{cr}$  & $M_{B,cr}$ & ${\\cal M}$ \\\\ \\hline 0.0~~~~~  & 1.573~~~~~ & 1.752~~~~~ & 1.573 \\\\ 1.0~~~~~  & 1.494~~~~~ & 1.643~~~~~ & 1.485 \\\\ 2.0~~~~~  & 1.390~~~~~ & 1.492~~~~~ & 1.361 \\\\ \\hline \\end{tabular} \\caption{Gravitational ($M_{cr}$) and baryonic ($M_{B,cr}$) critical mass (see text for more details) for proto-hadronic stars at different entropy per baryon  $\\tilde {S}/k_B$. ${\\cal M}$ denotes the gravitational mass of  the cold hadronic configuration with the same stellar baryonic mass ($M_{B,cr}$). Stellar masses are in units of the solar mass, $M_\\odot = 1.989 \\times 10^{33}$~g.} \\label{table} \\end{table}  In summary, in this work we have studied the quark deconfinement phase transition in hot $\\beta$-stable hadronic matter, and we have explored some of its consequences for the physics of neutron stars at birth. Our main finding is that proto-hadronic stars with a mass lower that the critical value $M_{cr}$ could survive the early stages of their evolution without decaying to a quark star. However, the prompt formation of a critical size drop of quark matter could take place when $M > M_{cr}$.  These proto-hadronic stars evolve to cold and deleptonized quark stars, or collapse to a black holes. Finally,  if quark matter nucleation occurs during post-bounce stage of core-collapse supernova, then the quark deconfinement phase transition could trigger a delayed supernova explosion characterised by a peculiar neutrino signal \\cite{sage+09}."
    },
    "0910/0910.3804_arXiv.txt": {
        "abstract": "{ We propose near-infrared colour selection criteria to extract Active Galactic Nuclei (AGNs) using the near-infrared Colour-Colour Diagram (CCD) and predict near-infrared colour evolution with respect to redshift. First, we cross-identify two AGN catalogues with the 2MASS Point Source Catalogue, and confirm both the loci of quasars/AGNs in the near-infrared CCD and redshift-colour relations. In the CCD, the loci of over $70 \\sim 80 \\%$ of AGNs can be distinguished from the stellar locus. To examine the colours of quasars, we simulate near-infrared colours using Hyperz code. Assuming a realistic quasar SED, we derive simulated near-infrared colours of quasars with redshift (up to $z \\sim 11$). The simulated colours can reproduce not only the redshift-colour relations but also the loci of quasars/AGNs in the near-infrared CCD. We finally discuss the possibility of contamination by other types of objects. We compare the locus of AGNs with the other four types of objects (namely, microquasars, cataclysmic variables, low mass X-ray binaries, and massive young stellar objects) which have a radiation mechanism similar to that of AGNs. In the near-infrared CCD, each type of object is located at a position similar to the stellar locus. Accordingly, it is highly probable that the four types of objects can be distinguished on the basis of the locus in a near-infrared CCD. We additionally consider contamination by distant normal galaxies. The near-infrared colours of several types of galaxies are also simulated using the Hyperz code. Although galaxies with $z\\sim 1$ have near-infrared colours similar to those of AGNs, these galaxies are unlikely to be detected because they are very faint. In other words, few galaxies should contaminate the locus of AGNs in the near-infrared CCD. Consequently, we can extract reliable AGN candidates on the basis of the near-infrared CCD. } ",
        "introduction": "Active Galactic Nuclei (AGNs) are tremendous energetic sources, where vast amounts of energy are generated by gravitational accretion around supermassive black hole. The radiation at nearly all wavelengths enables us to detect AGNs in multiwavelength observations. Hence, AGNs have been studied at various wavelengths. Past studies show that their Spectral Energy Distributions (SEDs) are roughly represented by a power-law (i.e., $f_\\nu \\propto \\nu^{-\\alpha}$), whilst normal galaxies produce an SED that peaks at $\\sim 1.6 \\mu$m as the composite blackbody spectra of the stellar population. Because the colours of an object provide us with rough but essential information about its spectrum, colours are important clues to identify AGNs from normal stars.  Colour selection is an efficient technique to distinguish AGNs from normal stars and have played an important role to extract AGN candidates without spectral observation. A classic method is known as the $UV$-excess \\citep[$UVX$; ][]{Sandage1965-ApJ,Schmidt1983-ApJ,Boyle1990-MNRAS}. The $UVX$ technique exploits the fact that quasars are relatively brighter than stars at shorter wavelength and therefore occupy a bluer locus in a CCD with respect to stars. In addition, many AGN candidates have been selected on the basis of colours in various wavelengths: optical \\citep{Richards2002-AJ}, optical and near-infrared \\citep{Glikman2007-ApJ}, and mid-infrared \\citep{Lacy2004-ApJS,Stern2005-ApJ}. These studies provide us with clues about the properties of AGNs.  Target selection of high redshift quasars has also been performed using their colours, mainly in optical wavelengths \\citep[e.g., ][]{Fan2000-AJ,Fan2001-AJ,Fan2003-AJ}. However, near-infrared properties are required when we try to select targets such as higher redshift quasars, since the shift of the Lyman break to longer wavelengths makes observations difficult at optical wavelengths. Therefore, near-infrared selection should be useful technique to extract high-redshift quasars.  In this paper, we present a study of the near-infrared colours of AGNs and demonstrate, by both observed and simulated colours, that the near-infrared colours can separate AGNs from normal stars. Additionally, we predict near-infrared colour evolution based on a Monte-Carlo simulation. In Sect. \\ref{Data}, we introduce the catalogues of AGNs which are used in order to investigate the observed colours. We confirm the near-infrared properties of spectroscopically confirmed AGNs on the basis of the near-infrared CCD and redshift-colour relations in Sect. \\ref{Properties}. In Sect. \\ref{Simulation}, we simulate the near-infrared colours using Hyperz code developed by \\citet{Bolzonella2000-AA} and demonstrate that AGNs reside in a distinct position in the near-infrared CCD. In Sect. \\ref{Discussion}, we consider the other probable objects which are expected to have near-infrared colours similar to those of AGNs.     \\if2 \\begin{figure*}[t] \\begin{tabular}{cc} (a) & (b) \\\\ \\begin{minipage}[tbp]{0.5\\textwidth} \\begin{center} \\resizebox{90mm}{!}{\\includegraphics[clip]{figure/qso-agn12-ccd.eps}} \\end{center} \\end{minipage} & \\begin{minipage}[htbp]{0.5\\textwidth} \\begin{center} \\resizebox{90mm}{!}{\\includegraphics[clip]{figure/sdss-qso-ccd.eps}} \\end{center} \\end{minipage} \\end{tabular} \\caption{(a) The distribution of AGNs in QA catalogue. (b) The distribution of quasars in SQ catalogue. The stellar locus in \\citet{Bessell1988-PASP} and the reddening vector taken from \\citet{Rieke1985-ApJ} are also shown. \\label{CCD1}} \\end{figure*} \\fi  \\begin{figure*}[tbp] \\begin{center} \\begin{tabular}{cc} (a) & (b) \\\\ \\resizebox{50mm}{!}{\\includegraphics[]{figure/qso_agn12-ccd-0z1.eps}} & \\resizebox{50mm}{!}{\\includegraphics[]{figure/sdss_qso-ccd-0z1.eps}} \\\\ \\resizebox{50mm}{!}{\\includegraphics[]{figure/qso_agn12-ccd-1z2.eps}} & \\resizebox{50mm}{!}{\\includegraphics[]{figure/sdss_qso-ccd-1z2.eps}} \\\\ \\resizebox{50mm}{!}{\\includegraphics[]{figure/qso_agn12-ccd-2z3.eps}} & \\resizebox{50mm}{!}{\\includegraphics[]{figure/sdss_qso-ccd-2z3.eps}} \\\\ \\resizebox{50mm}{!}{\\includegraphics[]{figure/qso_agn12-ccd-3z4.eps}} & \\resizebox{50mm}{!}{\\includegraphics[]{figure/sdss_qso-ccd-3z4.eps}} \\\\ \\resizebox{50mm}{!}{\\includegraphics[]{figure/qso_agn12-ccd-4z5.eps}} & \\resizebox{50mm}{!}{\\includegraphics[]{figure/sdss_qso-ccd-4z5.eps}} \\\\ \\end{tabular} \\caption{(a) The distribution of AGNs in the QA catalogue. (b) The distribution of quasars in the SQ catalogue. The stellar locus \\citep{Bessell1988-PASP}, the CTTS locus \\citep{Meyer1997-AJ}, and the reddening vector taken from \\citet{Rieke1985-ApJ} are also shown. \\label{CCD1}} \\end{center} \\end{figure*} ",
        "conclusions": "We confirmed the loci of catalogued quasars/AGNs in a ($H-K_\\textnormal{\\tiny S}$)-($J-H$) diagram, of which over $70 \\sim 80\\%$ are clearly separated from the stellar locus. In addition, we simulated the near-infrared colours of quasars on the basis of a Monte-Carlo simulation with Hyperz code, and demonstrated that the simulated colours can reproduce both the redshift-colour relations and the locus of quasars in the near-infrared CCD. We also predicted the colour evolution with respect to redshift (up to $z \\sim 11$). Finally, we discussed the possibility of contamination by other types of objects. The locus of AGNs is also different from those of the other four probable types of objects (namely, Microquasars, CVs, LMXBs, MYSOs) that are expected to be located at a similar locus. We also demonstrated by a Monte-Carlo simulation that normal galaxies are unlikely to contaminate the locus of AGNs in the near-infrared CCD.   \\citet{Hewett2006-MNRAS} investigated near-infrared colours of quasars using an artificial SED, but we proposed near-infrared colour selection criteria for extracting AGNs and studied both observed and simulated colours with quantitative discussions. An important point is that our selection criteria require only near-infrared photometric data, although some previous studies \\citep[e.g., ][]{Glikman2007-ApJ,Glikman2008-AJ} used colour selections based on colours between near-infrared and optical wavelengths. In other words, our selection criteria make extraction of candidates easier because only near-infrared colours are needed. This technique should also be useful when we search for high-redshift quasars, since they become very faint in the optical wavelength due to the shift of Lyman break.   This paper demonstrates that near-infrared colours can be useful to select AGN candidates. If an additional constraint is imposed, more reliable candidates can be extracted. When we use the near-infrared colour selection with an additional constraint for near-infrared catalogues containing sources distributed in large area (e.g., 2MASS, DENIS, UKIDSS, and future surveys), a lot of AGN samples (possibly over $\\sim$10 000) are expected to be derived in a region over $\\sim$10 000 deg$^2$. \\citet{Kouzuma2009-prep} \\citep[see also][]{Kouzuma2009-ASPC} cross-identified between the 2MASS and ROSAT catalogues, and extracted AGN candidates in the entire sky using the near-infrared colour selection in this paper. These large number of samples may provide us with clues about such as the evolution of AGNs and X-ray background. Additionally, in our simulation, quasars with $z \\gtrsim 8$ can be extracted on the basis of near-infrared colours. This property might be helpful to search for high-redshift quasars in the future."
    },
    "0910/0910.4279_arXiv.txt": {
        "abstract": "The Herschel ATLAS is the largest open-time key project that will be carried out on the Herschel Space Observatory. It will survey 510 square degrees of the extragalactic sky, four times larger than all the other Herschel surveys combined, in five far-infrared and submillimetre bands. We describe the survey, the complementary multi-wavelength datasets that will be combined with the Herschel data, and the six major science programmes we are undertaking. Using new models based on a previous submillimetre survey of galaxies, we present predictions of the properties of the ATLAS sources in other wavebands. ",
        "introduction": "Approximately half the energy emitted since the big bang by all the objects in the Universe has been absorbed by dust and then reradiated between 60 and 500 $\\mu$m \\citep{dwek98,fix98,driver1}, a wavelength range in which the Universe is still largely unexplored (Fig. 1). On the short-wavelength side of this waveband, the whole sky was surveyed at 60 and 100 $\\mu$m by IRAS in the 1980s. However, almost all of the tens of thousands of galaxies detected by IRAS were spirals and starbursts in the nearby Universe ($\\rm z < 0.1$), and IRAS revealed little about the dust in other galaxy populations, especially early-type galaxies \\citep{breg98}. Even in the late-type galaxies, only the small fraction of the the dust warm enough to radiate significantly in the far-infrared was detected by IRAS. Devereux and Young (1990), for example, showed that the gas-to-dust ratio estimated from IRAS measurements alone was $\\simeq$10 times greater than the standard Galactic value, implying that $\\simeq$90\\% of the dust in galaxies was effectively missed by IRAS. ISO and Spitzer, with their long wavelength (170 $\\mu$m) band, suffered less from this problem but will still have effectively missed any dust with $\\rm T < 15 K$ \\cite{ben2003}. Notwithstanding the successes of IRAS, ISO and Spitzer, most of the waveband from 60 to 500 $\\mu$m is still virtually {\\it terra incognita}, and the only survey of a large area of the extragalactic sky at a wavelength beyond 200 $\\mu$m is the one recently carried out by the Herschel pathfinder experiment, the Balloon Large Area Survey Telescope (BLAST), which covers $\\simeq$ 20 deg$^2$ \\citep{devlin}.  The lack of a survey covering a large area of sky in the submillimetre waveband (in this paper defined as $\\rm 100\\ \\mu m < \\lambda < 1\\ mm$) has left us in some ways knowing more about dust in the distant, early Universe than in the Universe today. The surveys that have been carried out in the submillimetre waveband with ground-based telescopes---at 450 and 850 $\\mu$m with the SCUBA camera on the James Clerk Maxwell Telescope \\cite{h98,eales99,cop2006}, at 1.2 mm with MAMBO on the IRAM 30-m telescope \\cite{greve2004,bert2007,greve2008} and at 1.1 mm with AZTEC on the James Clerk Maxwell Telescope \\cite{per08,scott08}---have been of very small areas of sky, covering $\\rm \\sim 1\\ deg^2$ of sky in total. Because of the unusual submillimetre `K-correction'\\footnote{Beyond a redshift of $\\simeq1$, as the redshift increases, the typical spectral energy distribution of a dusty galaxy means that the effect of increasing luminosity-distance on the brightness of the galaxy is compensated for by the increasing rest-frame luminosity of the galaxy. The consequence is that the galaxy's flux density is approximately independent of redshift.}, these surveys have mostly detected sources at very high redshifts ($\\rm z \\succeq 1$). They have led to the important discovery that there is a population of luminous dust-enshrouded galaxies in the early Universe \\cite{sib,h98}, which many authors have suggested are the ancestors of ellipticals today \\cite{scott02,dunne2003}, but they have told us relatively little about the evolution of the Universe since $\\rm z \\simeq 1$ (the last 8 billion years) and almost nothing at all about dusty galaxies in the nearby Universe.  \\begin{figure} \\figurenum{1} \\plotone{Fig1.eps} \\caption{The extraglactic background radiation as a function of wavelength \\citep{dole}} \\end{figure}  Two basic things one would like to know about the nearby Universe are the submillimetre luminosity function and the dust-mass function(the space-density of galaxies as a function of dust mass). These functions are important for many reasons, including tests of semi-analytical models of galaxy formation \\cite{cole,baugh} and, by comparison with the same functions at high redshift,  accurate measurements of the amount of evolution that is occurring in the submillimetre waveband \\citep{dunne2003}. Unfortunately, our knowledge of these functions is still extremely poor because of the limitations in areal coverage and sensitivity of previous submillimetre telescopes. Until recently the only estimates of the local submillimetre luminosity function were one based on 55 galaxies detected in an ISO 170-$\\mu$m survey \\citep{tak2006} and ones based on SCUBA 850-$\\mu$m observations of $\\simeq$200 galaxies selected in other wavebands \\citep{dunne2000,cat2005}. There are now direct estimates from the BLAST results of the local luminosity function at 250, 350 and 500 $\\mu$m \\citep{eales09}, but their accuracy is limited by the small number of low-redshift sources detected in the BLAST survey: $\\rm \\simeq 30$ at $\\rm z < 0.2$.  The lack of a large-area survey capable of measuring the dust content and dust-obscured star formation in large numbers of galaxies in the local Universe has been especially galling for submillimetre astronomers in the light of the success of their colleagues working in other wavebands. The Sloan Digital Sky Survey and the 2dF Galaxy Redshift Survey have led to a revolution in our understanding of the distribution of galaxies in the local Universe, and the relationships between their present star-formation rate, star-formation history, stellar mass, morphology and environment \\cite{lew02,kauff03a,kauff03b,alan04,balogh04}. However, all the studies that have used these impressive datasets to investigate the physics and ecology of the galaxy population have been forced to ignore the dust phase of the interstellar medium and star formation that is heavily obscured by dust, because the IRAS survey was only sensitive enough to detect a small percentage of the galaxies in the redshift surveys (1.8 per cent of the SDSS galaxies - Obric et al. 2006), and it missed 90\\% of the dust in the detected galaxies because of its insensitivity to cold dust \\citep{dy90}.  The launch of the Herschel Space Observatory, which occurred on May 14th 2009, has the potential to dramatically increase our knowledge of dust and dust-obscured star formation, especially in the nearby Universe. Herschel has two main cameras: SPIRE, which will be able to image the sky simultaneously at 250, 350 and 500 $\\mu$m \\cite{matt}, and PACS, which will be able to image the sky in two bands simultaneously, either 70 and 170 $\\mu$m or 110 and 170 $\\mu$m \\cite{pog}. Herschel will have much better angular resolution ($\\simeq$18 arcsec at 250 $\\mu$m) and sensitivity than previous observatories, and the spectral coverage of SPIRE will make it possible to carry out the first large-area surveys in this virtually unexplored part of the electromagnetic spectrum. The SPIRE bands will also make it possible to detect the cold dust that was missed by earlier observatories. Herschel is also better suited for investigating the local Universe than the submillimetre surveys that will soon be carried out from the ground, in particular the SCUBA-2 and LABOCA surveys, because these will operate mostly at 850 $\\mu$m, where low-redshift galaxies are intrinsically faint, whereas the Herschel bands span the peak of the typical spectral energy distribution of a galaxy in the nearby Universe.  In this paper, we describe the largest project that will be carried out with the Herschel Space Observatory in `Open Time', the time available for competition within the international astronomical community\\footnote{This forms about 2/3 of the total observing time on Herschel, with the remaining 1/3 being time reserved for the teams that built the instruments---`Guaranteed Time'.}. The Herschel Astrohysical Terahertz Large Area Survey (the Herschel ATLAS or H-ATLAS) will be a survey of 510 deg$^2$ of sky in five photometric bands. This is $\\simeq$8 times larger than the coverage of the next largest (in area) Herschel extragalactic survey, HERMES (Oliver et al. 2009). The main scientific goal of the H-ATLAS is to provide measurements of the dust masses and dust-obscured star formation for tens of thousands of nearby galaxies, the far-IR/submillimetre equivalent to the SDSS photometric survey. However, the H-ATLAS has many other science goals ranging from the investigation of the point sources that will be detected by the Planck Surveyor to a study of high-latitude galactic dust.  The arrangement of this paper is as follows. In Section 2 we describe the basic parameters of the survey. In Section 3 we present predictions of the number and redshift distribution of the sources that will be detected by the H-ATLAS. In Section 4, we describe the complementary data that exists or will soon exist for the H-ATLAS fields. In Section 5 we describe the six main H-ATLAS science programmes. In Section 6 we describe the detailed survey strategy, including issues that will be of interest to the general community, such as our plans for the release of data products. We everywhere assume the cosmological parameters for a `concordance universe': $\\rm \\Omega_M = 0.267$ and $\\rm \\Omega_{\\Lambda} = 0.762$.  \\begin{figure*} \\figurenum{2} \\plottwo{Figure2a.ps}{Figure2b.ps} \\caption{The positions of the ATLAS field, shown as white blocks, superimposed on the IRAS 100 $\\mu$m map of the sky, which traces the distribution of galactic dust. Figure 2(a) shows the northern galactic cap and Figure 2(b) shows the southern galactic cap. The colour coding for the lines is as follow: solid green lines---RA and dec; dotted green lines---ecliptic latitude and longitude; cyan---KIDS/VIKING area; yellow---SDSS area; blue---2dFGRS area; magenta---area of the Dark Energy Survey; magenta/blue dashed---area covered by the South Pole Telescope.} \\end{figure*} ",
        "conclusions": "The Herschel ATLAS is the largest open-time key project that will be carried out with the Herschel Space Observatory. It will survey 510 square degrees of the extragalactic sky, four times larger than all the other Herschel surveys combined, in five far-infrared and submillimetre bands. We have described the survey, the complementary multi-wavelength datasets that will be combined with the Herschel data, and the six major science programmes we plan to undertake. Using new models based on a submillimetre survey of nearby galaxies, we have presented predictions of the properties of the sources that will be detected by the H-ATLAS. We intend to release the H-ATLAS data---maps and catalogues---to the astronomical community in a series of data releases. The first of these will be in May 2010."
    },
    "0910/0910.0675.txt": {
        "abstract": "{ We have collected spectra of about 2000 red giant branch (RGB) stars in 19 Galactic globular clusters (GC) using FLAMES@VLT (about 100 star with GIRAFFE and about 10 with UVES, respectively, in each GC).  These observations provide an unprecedented, precise, and homogeneous  data-set of Fe abundances in GCs. We use it to study the cosmic scatter of iron and find that, as far as Fe is concerned, most GCs can still be considered mono-metallic, since the upper limit to the scatter in iron is less than 0.05 dex, meaning that the degree of homogeneity is better than 12\\%. The scatter in Fe we find seems to have a dependence on luminosity, possibly due to the well-known inadequacies of stellar atmospheres for upper-RGB stars and/or to intrinsic variability. It also seems to be correlated with cluster properties, like the mass, indicating a larger scatter in more massive GCs which is likely a (small) true intrinsic scatter. The 19 GCs, covering the metallicity range of the bulk of Galactic GCs, define an accurate and updated  metallicity scale. We provide transformation equations for a few existing scales. We also provide new values of [Fe/H], on our scale, for all GCs in the Harris' catalogue. } ",
        "introduction": "Nowadays, the chemical composition of globular cluster (GC) stars cannot be regarded anymore as strictly homogeneous, as implied from the classical paradigm that these old stellar aggregates were the best approximation in nature of simple stellar populations (see Gratton, Sneden and Carretta 2004 for an extensive review and references). Recent investigations using large samples of stars observed with multi-object spectrographs at large telescopes highlight that every GC studied so far harbours at least two different stellar generations, distinct in chemical composition and age (Carretta et al. 2009a, 2009b).  These different populations are found to differ in abundances of light elements (C, N, O,  Na, Mg, Al, Si,  F; Smith and Martell 2003; Smith et al. 2005; Carretta et al. 2009a,b; Yong et al. 2008a,b; Melendez and Cohen 2009, to quote a few studies; see also  Gratton et al. 2004 and references therein) involved in proton-capture reactions of H-burning at high temperature (Denisenkov and Denisenkova 1989; Langer et al. 1993). Star to star abundance variations in these elements are expected to also come with differences in the main outcome of the H burning, the He content (Gratton et al. 2009;  Bragaglia et al. in preparation; Prantzos and Charbonnel 2006; Ventura et al. 2001). Apart from a few alterations presently well understood in term of an extra-mixing episode after the red giant branch (RGB) bump (Charbonnel 1994, 1995; Charbonnel and Zahn 2007;  Eggleton et al. 2007) the abundance variations are inherited by currently observed, long  lived GC stars from a previous stellar component/generation: both spectroscopic (Gratton et al. 2001; Ramirez and Cohen 2002; Carretta et al. 2004; Piotto et al. 2005) and photometric  (e.g. Bedin et al. 2004) observations convincingly showed that the observed pattern of chemical composition is present also among unevolved stars on the subgiant branch and the main sequence.  On the other hand, apart from a few notable exceptions\\footnote{The GC $\\omega$ Cen (widely considered to be the remnant of a dwarf galaxy) shows evidences of several bursts of star formations, with corresponding peaks in the metal abundance (see Gratton et al. 2004 for a comprehensive discussion) and a spread in metallicity seems to be confirmed in M~22  (Marino et al. 2009) as well as in M~54 (Bellazzini et al. 2008 and Carretta et al., in preparation).} GCs are still found mono-metallic objects, as far as abundances of heavier elements are concerned (see Gratton et al. 2004 for a recent review on this subject). Their heavy (Z$>$13) elements metallicity, usually represented by the ratio [Fe/H]\\footnote{We adopt the usual spectroscopic notation, $i.e.$ [X]= log(X)$_{\\rm star} -$ log(X)$_\\odot$ for any abundance quantity X, and log $\\epsilon$(X) = log (N$_{\\rm X}$/N$_{\\rm H}$) + 12.0 for absolute number density abundances.}, is found to be extremely homogeneous from star to star in each cluster.  The sites of production of heavy elements (in particular $\\alpha-$capture elements, Fe-group elements) are stars with large and intermediate initial masses, exploding as core-collapse or thermonuclear supernovae (see Wheeler, Sneden and Truran 1989). By studying the level and the dispersion of their yields inherited by stars in GCs we have the possibility of investigating the past history of the precursors from whom the present-day globular clusters formed (see Carretta et al. 2009c).  The recently completed analysis of high resolution spectra of almost 2,000 RGB stars in  19 Galactic GCs (Carretta et al. 2009a,b and references therein) provides an unprecedented sample of  stars analysed in a fully homogeneous way. By exploiting these data we can examine in detail the issue of the cosmic scatter in the metallicity (hereafter the abundance of Fe-peak elements) of GCs and hopefully provide clues to the early evolution of their progenitors.  Moreover, abundances of iron from high resolution spectra are traditionally the calibrating points for metallicity scales based on photometric and/or low-resolution spectroscopic indices that are sensitive to metal abundances (see the excellent discussion and historical review on this issue by Kraft and Ivans 2003).  Thus, our sample is perfectly suited to provide robust (for statistics) and homogeneous  (for technique of analysis) \"pillars\" onto which anchor several existing metallicity scales.  The paper is organized as follows. We give in section 2 a brief description of the data set and of the analysis. We discuss the intrinsic (cosmic) spread of iron in section 3, where we also present some correlations of the spread with GC properties, like the absolute total luminosity or the $\\alpha-$elements abundance. We discuss the dependence of this spread on the stars luminosity in section 4 and present a recalibration of four metallicity scales to our new one in section 5. A summary is given in section 6 and we give metallicity values on our new scale for all GCs in the Harris (1996) catalogue in the Appendix. ",
        "conclusions": "Within our survey of 19 Galactic GCs with FLAMES (UVES and GIRAFFE), we have measured the iron abundance of about 2000 RGB stars in a accurate and homogeneous way (determination of atmospheric parameters, measure of $EW$s, spectroscopic analysis, etc). In the present paper we  exploit this huge, homogeneous data set, covering (almost) the entire range of GC metallicity in our Galaxy and a summary of our findings follows: \\begin{itemize} \\item [i)] we obtain a measure of the intrinsic spread of iron in GCs from the $rms$ scatter of our metallicities. We confirm that the cosmic scatter in Fe of most GCs is very small: the upper limit to this scatter is less than 0.05 dex, i.e., the iron abundance is homogeneous within 12\\% in each GC, on average; \\item [ii)] the $rms$ scatter we find for the UVES sample within each GC is larger than the one derived for the GIRAFFE sample. This can be ascribed to the different magnitude ranges observed. The UVES stars are usually brighter, nearer the RGB tip, and may suffer from larger inadequacies in the treatment  of atmospheres and intrinsic stellar variability; \\item[iii)] this is confirmed by a increase in scatter for brighter stars also in the GIRAFFE samples, that is not explained by AGB contamination, and that is larger for metal-poor GCs; \\item[iv)] there are interesting correlations between the $rms$ scatter in Fe and several GC parameters, like $M_V$, $T_{eff}^{max}$(HB), $\\alpha$-element abundance, that seem to indicate a better capability of more massive clusters to self-enrich; \\item [v)] we recalibrate several metallicity scales of GCs using our 19 precise, homogeneous values; we have done it for the GC97, ZW, KI03, and R97 ones, giving transformation equations; \\item [vi)] finally, in the Appendix we give [Fe/H] on our new scale for all clusters in the Harris (1996) compilation. \\end{itemize}"
    },
    "0910/0910.2175_arXiv.txt": {
        "abstract": "We investigate the effects of the nonminimal coupling between the scalar field dark energy (quintessence) and the dark matter on the two-point correlation function. It is well known that this coupling shifts the turnover scale as well as suppresses the amplitude of the matter power spectrum. However, these effects are too small to be observed when we limit the coupling strength to be consistent with observations. Since the coupling of quintessence to baryons is strongly constrained, species dependent coupling may arise. This results in a baryon bias that is different from unity. Thus, we look over the correlation function in this coupled model. We are able to observe the enhancement of the baryon acoustic oscillation (BAO) peak due to the increasing bias factor of baryon from this species dependent coupling. In order to avoid the damping effect of the BAO signature in the matter power spectrum due to nonlinear clustering, we consider the coupling effect on the BAO bump in the linear regime. This provides an alternative method to constrain the coupling of dark energy to dark matter. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.2248.txt": {
        "abstract": "{} {The aim of this paper is to discuss the nature of two Type Ic supernovae SN 2007bg and SN 2007bi and their host galaxies. Both supernovae were discovered in wide-field, non-targeted surveys and are found to be associated with sub-luminous blue dwarf galaxies identified in SDSS images.} {We present $BVRI$ photometry and optical spectroscopy of SN 2007bg and SN 2007bi and their host galaxies. Their lightcurves and spectra are compared to those of other Type Ic SNe and analysis of these data provides estimates of the energetics, total ejected masses and synthesised mass of $^{56}$Ni.  Detection of the host galaxy emission lines allows for metallicity measurements.} {Neither SNe 2007bg nor 2007bi were found in association with an observed GRB, but from estimates of the metallicities of their host-galaxies they are found to inhabit similar low-metallicity environments as GRB associated supernovae. The radio-bright SN 2007bg is hosted by an extremely sub-luminous galaxy of magnitude $M_{B}=-12.4${\\tiny$\\pm0.6$} mag and an estimated oxygen abundance of $12+\\log({\\rm O/H})=8.18${\\tiny $\\pm0.17$} (on the \\citealt{Pettini:2004p217} scale). The early lightcurve evolution of SN 2007bg matches the fast-pace decline of SN 1994I giving it one of the fastest post-maximum decline rates of all broad-lined Type Ic supernovae known to date and, when combined with its high expansion velocities, a high kinetic energy to ejected mass ratio ($E_K/M_{ej}\\sim2.7$). We also show that SN 2007bi is possibly the most luminous Type Ic known, reaching a peak magnitude of $M_{R} \\sim -21.3$ mag and  displays a remarkably slow decline, following the radioactive decay rate of $^{56}$Co to $^{56}$Fe throughout the course of its observed lifetime. SN 2007bi also displays an extreme longevity in its spectral evolution and is still not fully nebular at approximately one year post-maximum brightness. From a simple model of the bolometric light curve of SN 2007bi we estimate a total ejected $^{56}$Ni mass of $M_{Ni}=3.5-4.5{\\rm M}_\\odot$, the largest $^{56}$Ni mass measured in the ejecta of a supernova to date. There are two models that could explain the high luminosity and large ejected $^{56}$Ni mass. One is a pair-instability supernova (PISN) which has been predicted to occur for massive stars at low metallicities. We measure the host galaxy metallicity of SN 2007bi to be $12+\\log({\\rm O/H}) = 8.15${\\tiny $\\pm0.15$} (on the \\citealt{McGaugh:1991p7985} scale) which is somewhat high to be consistent with the PISN model. An alternative is the core-collapse of a C+O star of $20-40$M$_{\\odot}$ which is the core of a star of originally $50-100$M$_{\\odot}$.} {} ",
        "introduction": "The absence of obvious Hydrogen and Helium in the spectra of supernovae Type Ic (SNe Ic) indicates that these objects originate from the gravitational core-collapse of stars that have been stripped of their outer layers during the course of their evolution. This stripping is thought to occur either via binary interaction or mass-loss resulting from the strong stellar winds suffered during the life-time of the progenitor star \\cite{Smartt:2009p11794}. As a class SNe Ic are known to exhibit a broad range in both luminosity and light-curve shape \\citep{Richardson:2006p72}; a heterogeneity also displayed in their spectra \\citep{Matheson:2001p4883,Elmhamdi:2006p6842}. Interest in this class of SN has been rejuvenated in the past decade with confirmed association of SNe Ic with long-duration Gamma-Ray Bursts (GRBs). In every case that a SN has been identified in association with a GRB it has been classified as a broad-lined SN Ic (SN Ic-BL): GRB 980425/SN 1998bw \\citep{Galama:1998p4661}, GRB 030329/SN 2003dh \\citep{Stanek:2003p8187, Hjorth:2003p8194, Matheson:2003p8201}, GRB 031203/SN 2003lw \\citep{Malesani:2004p8254, Mazzali:2006p8264} and GRB 060218/SN 2006aj \\citep{Modjaz:2006p4901, Pian:2006p8257, Mazzali:2006p8258, Campana:2006p8259, Sollerman:2006p4932, Mirabal:2006p8260, Cobb:2006p8263}. The broad-lines exhibited by these SNe are indicative of the huge photospheric velocities achieved in their explosions.  \\begin{figure*}\t\t%%%%%%%%%%%%  2 X SEQUENCE STAR FINDER IMAGES %%%%%%%%%%%%%%%%% \\begin{minipage}[t]{0.48\\linewidth}\t\t%SN2007bg% \\centering \\includegraphics[scale=0.52]{./figures/13004f01.eps} \\caption{The field of SN 2007bg. Local standard stars used to calibrate the photometry of the SN are labelled $1-6$. TNG {\\it B}-band image taken on 2007 May 11th, 25 days post-discovery.} \\label{fig:2007bg_sequence_fig} \\end{minipage} \\hspace{0.5cm} \\begin{minipage}[t]{0.48\\linewidth}\t\t%SN2007bi% \\centering \\includegraphics[scale=0.51]{./figures/13004f02.eps} \\caption{The field of SN 2007bi. Local standard stars used to calibrate the photometry of the SN are labelled $1-9$. LT {\\it R}-band image taken on 2007 May 12th, 36 days post-discovery. } \\label{fig:2007bi_sequence_fig} \\end{minipage} \\end{figure*}  \\cite{Modjaz:2008p87} measured the metallicities at the sites of all GRB associated SNe Ic-BL (GRB/SNe Ic-BL) and compared them to the metallicities measured at the sites of those SNe Ic-BL not found in association with a GRB (non-GRB/SNe Ic-BL). Regardless of the abundance scale employed, \\citeauthor{Modjaz:2008p87} found that the environment of every non-GRB/SN Ic-BL was more metal-rich than the environment of any of the GRB/SNe Ic-BL. Upon studying the host galaxies of 42 GRBs discovered within the {\\it Hubble} Higher-z Supernova Search, \\cite{Fruchter:2006p234} found that GRBs ({\\it a}) typically favour low-luminosity, blue, irregular galaxies and ({\\it b}) tend to concentrate on the very brightest regions of these host galaxies. This evidence adds weight to the theory that the progenitors of GRBs are massive stars that inhabit low-metallicity environments \\citep[e.g. the `Collapsar' model: ][]{Woosley:1993p8292,Woosley:2006p167}.  With the advent of wide-field, non-targeted SN surveys, more SNe are now being detected in sub-luminous hosts and many of these SNe are proving to be peculiar, e.g. the over-luminous SNe 2005ap \\citep{Quimby:2007p50} and 2008es \\citep{Gezari:2009p2983,Miller:2009p2870}. Here we report on the follow-up of two such objects discovered by non-targeted SN surveys, both hosted by low-luminosity, dwarf galaxies: SNe 2007bg and 2007bi.  SN 2007bg was discovered in an unfiltered image taken with the 0.45m ROTSE-IIIb telescope \\citep[Robotic Optical Transient Search Experiment,][]{Akerlof:2003p26} at the McDonald Observatory, Texas by the Texas Supernova Survey. The object was discovered on 2007 April 16.15 (UT dates are used throughout this paper) located at $\\alpha = 11^{h}49^{m}26^{s}.18$, $\\delta = +51^{o}49^{'}21^{''}.8$ (J2000.0) \\citep{Quimby:2007p3598}. The SN was also recovered in unfiltered CCD images taken by the ROTSE-IIId telescope in Bakirlitepe, Turkey on 2007 April 16.81. No source was recovered at the location of the SN in images taken by the ROTSE-IIIb telescope on 2007 April 6.15, providing an upper-limit to the magnitude of the SN at that epoch. A spectrum taken on 2007 April 18.30 with the 9.2-m Hobby-Eberly Telescope led \\citeauthor{Quimby:2007p3598} to classify this object as a Type Ic, claiming similarity with SN 2002ap at around 11 days post maximum brightness. Upon obtaining a spectrum of SN 2007bg with the Telescopio Nazionale Galileo (TNG) on 2007 April 21.04, \\citet{Harutyunyan:2007p5835} confirm the resemblance to other SNe Ic-BL. They report that the spectrum is dominated by a broad, strong feature centered at about 6524$\\AA$ and state that if this feature is associated with H$\\alpha$ at a velocity of $\\sim$18 000 km s$^{-1}$ then SN 2007bg should be re-classified as a broad-lined Type II similar to SN 2003bg \\citep{Hamuy:2003p6874,Hamuy:2009p12161,Mazzali:2009p12162}.  \\begin{table}[t]\t\t%%%%%%%%%%%%  SEQUENCE STAR TABLE FOR SN 2007bg %%%%%%%%%%%%%%%% \\begin{center} \\caption{\\textrm {Local standard stars used to calibrate the photometry of SN 2007bg to the {\\it BVRI} standard Landolt system (Johnson-Cousins).}\\label{2007bg_sequence}} \\begin{tabular}{c c c c c} \\hline\\hline {Star ID} & {\\it B} & {\\it V} & {\\it R} & {\\it I}  \\\\ \\hline 1\t& 19.45  {\\tiny$\\pm$\t0.03}\t\t& 19.07  {\\tiny$\\pm$\t0.04}\t\t& 18.72  {\\tiny$\\pm$ 0.03}\t\t& 18.53  {\\tiny$\\pm$\t0.04}\t \\\\ 2\t& 20.71  {\\tiny$\\pm$\t0.07}\t\t& 19.36  {\\tiny$\\pm$\t0.04}\t\t& 18.42  {\\tiny$\\pm$ 0.05}\t\t& 17.06  {\\tiny$\\pm$\t0.02}\t \\\\ 3\t& 17.85  {\\tiny$\\pm$\t0.02}\t\t& 17.29  {\\tiny$\\pm$\t0.02}\t\t& 16.87  {\\tiny$\\pm$ 0.02}\t\t& 16.58  {\\tiny$\\pm$\t0.02}\t \\\\ 4\t& 18.91  {\\tiny$\\pm$\t0.03}\t\t& 18.03  {\\tiny$\\pm$\t0.02}\t\t& 17.34  {\\tiny$\\pm$ 0.03}\t\t& 16.86  {\\tiny$\\pm$\t0.02}\t\\\\ 5\t& 19.68  {\\tiny$\\pm$\t0.04}\t\t& 19.34  {\\tiny$\\pm$\t0.04}\t\t& 19.08  {\\tiny$\\pm$ 0.04}\t\t& 18.82  {\\tiny$\\pm$\t0.04}\t \\\\ 6\t& 16.00  {\\tiny$\\pm$\t0.02}\t\t& 15.49  {\\tiny$\\pm$\t0.02}\t\t& 15.02  {\\tiny$\\pm$ 0.02}\t\t& 14.72  {\\tiny$\\pm$\t0.02}\t \\\\ \\hline \\end{tabular} \\end{center} \\end{table}  SN 2007bi was discovered in images taken by the QUEST II camera on the Palomar Oschin 1.2-m Schmidt telescope as part of the NEAT component of the Palomar-QUEST Consortium \\citep{Djorgovski:2008p1193} by the Nearby Supernova Factory collaboration \\citep{Aldering:2002p3} on 2007 April 6.5 \\citep{Gal-Yam:2009}. The object is located at $\\alpha = 13^{h}19^{m}20^{s}.19$, $\\delta = +08^{o}55^{'}44^{''}.3$ (J2000.0). A spectrum of the object obtained with the Supernova Integral Field Spectrograph on the University of Hawaii 2.2-m telescope under poor conditions on 2007 April 8.6 led \\citeauthor{Gal-Yam:2009} to tentatively classify SN 2007bi as a Type Ic. Further spectra taken on 2007 April 15.6 and April 16.4 with the Low Resolution Imaging Spectrometer on the Keck I 10-m telescope, again under poor conditions, helped \\citeauthor{Gal-Yam:2009} to confirm the classification of SN 2007bi as a peculiar Type Ic, analogous to SN 1999as \\citep{Knop:1999p6878, Hatano:2001p4590} at 25 days post-discovery.  \\begin{table}[t]\t\t%%%%%%%%%%%%  SEQUENCE STAR TABLE FOR SN 2007bi %%%%%%%%%%%%%%%% \\begin{center} \\caption{\\textrm {Local standard stars used to calibrate the photometry of SN 2007bi to the {\\it BVRI} standard Landolt system (Johnson-Cousins).}\\label{2007bi_sequence}} \\begin{tabular}{c c c c c} \\hline\\hline {Star ID} & {\\it B} & {\\it V} & {\\it R} & {\\it I}  \\\\ \\hline 1\t& 17.91  {\\tiny$\\pm$\t0.02}\t\t& 16.53  {\\tiny$\\pm$\t0.02}\t\t& 15.60  {\\tiny$\\pm$ 0.02}\t\t& 14.78  {\\tiny$\\pm$\t0.01}\t \\\\ 2\t& 18.27  {\\tiny$\\pm$\t0.02}\t\t& 17.82  {\\tiny$\\pm$\t0.01}\t\t& 17.53  {\\tiny$\\pm$ 0.02}\t\t& 17.17  {\\tiny$\\pm$\t0.01}\t \\\\ 3\t& 18.38  {\\tiny$\\pm$\t0.02}\t\t& 17.85  {\\tiny$\\pm$\t0.02}\t\t& 17.52  {\\tiny$\\pm$ 0.02}\t\t& 17.13  {\\tiny$\\pm$\t0.01}\t \\\\ 4\t& 17.15  {\\tiny$\\pm$\t0.01}\t\t& 16.45  {\\tiny$\\pm$\t0.01}\t\t& 15.99  {\\tiny$\\pm$ 0.01}\t\t& 15.58  {\\tiny$\\pm$\t0.01}\t\\\\ 5\t& 17.17  {\\tiny$\\pm$\t0.01}\t\t& 16.65  {\\tiny$\\pm$\t0.01}\t\t& 16.31  {\\tiny$\\pm$ 0.01}\t\t& 15.95  {\\tiny$\\pm$\t0.01}\t \\\\ 6\t& 18.56  {\\tiny$\\pm$\t0.02}\t\t& 18.00  {\\tiny$\\pm$\t0.02}\t\t& 17.68  {\\tiny$\\pm$ 0.02}\t\t& 17.32  {\\tiny$\\pm$\t0.01}\t \\\\ 7\t& 19.92  {\\tiny$\\pm$\t0.03}\t\t& 18.42  {\\tiny$\\pm$\t0.02}\t\t& 17.50  {\\tiny$\\pm$ 0.02}\t\t& 16.56  {\\tiny$\\pm$\t0.01}\t \\\\ 8\t& 17.08  {\\tiny$\\pm$\t0.02}\t\t& 16.05  {\\tiny$\\pm$\t0.01}\t\t& 15.32  {\\tiny$\\pm$ 0.01}\t\t& 14.78  {\\tiny$\\pm$\t0.01}\t \\\\ 9\t& 19.12  {\\tiny$\\pm$\t0.03}\t\t& 17.81  {\\tiny$\\pm$\t0.02}\t\t& 16.89  {\\tiny$\\pm$ 0.02}\t\t& 16.18  {\\tiny$\\pm$\t0.01}\t \\\\ \\hline \\end{tabular} \\end{center} \\end{table}  In this paper we present {\\it BVRI} photometric and spectroscopic data sets for SNe 2007bg and 2007bi, discussing techniques of data reduction and our estimations of the distances and extinction towards both SNe in Sections \\ref{phot_obs} and \\ref{redshifts}. In Sections \\ref{lightcurves} and \\ref{bol_lcs} we analyse the lightcurves of both SNe, comparing them with other SNe Ic. In Section \\ref{spectra} we present spectroscopic data for both SNe, describing the techniques of data reduction. Sections \\ref{spec_evol} and \\ref{spec_comp} are dedicated to the analysis of the spectra and their comparison with other SNe. In Section \\ref{other_para} we present the photospheric velocities, a derivation of the explosion parameters, the search for associated GRBs and an estimation of the oxygen abundances of the host galaxies. Finally we discuss and conclude our findings in Sections \\ref{Discussion} and \\ref{Conclusions}.  \\begin{table*}\t\t%%%%%%%%%%%%%%%  TABLE OF PHOTOMETRY FOR SN 2007bg  %%%%%%%%%%%%%% \\setlength{\\tabcolsep}{5pt} \\begin{minipage}[t]{1.0\\textwidth} \\begin{center} \\caption{\\textrm {Photometry of SN 2007bg.}\\label{2007bg_phot}} \\begin{tabular}{l r r r c r c r c r c r l} \\hline\\hline {UT Date}\t& {JD}\t\t& Epoch$^{a}$\t& Rest-frame\t\t&{\\it B}\t& {$K_{BB}$}$^{b}$\t&{\\it V}\t& {$K_{VV}$}$^{b}$\t& {\\it R}\t& {$K_{RR}$}$^{b}$\t& {\\it I}\t& {$K_{II}$}$^{b}$\t& {Source}$^{c}$\t\\\\ & {- 2,450,000}\t& (d)\t\t\t& epoch (d)$^{a}$\t&\t\t&  \t\t\t\t&\t\t&\t\t\t\t&\t\t&\t\t\t\t&\t\t&\t\t\t\t&\t\t\t\t\\\\ \\hline 06/04/07\t& 4196.7\t& $-$10.0\t& $-$9.6\t&\t\t\t\t\t&\t\t&\t\t\t\t\t&\t\t&\t\\multicolumn{1}{l}{(18.20)$^{d}$}\t&\t\t&\t\t\t\t\t\t&\t\t&\tR-IIIb$^{e}$ \\\\ 16/04/07\t& 4206.7\t& 0.0\t\t& 0.0\t\t&\t\t\t\t\t&\t\t&\t\t\t\t\t&\t\t&\t\\multicolumn{1}{l}{17.7}\t\t\t& $-$.05\t&\t\t\t\t\t\t&\t\t&\tR-IIIb$^{e}$ \\\\ 16/04/07\t& 4207.3\t& 0.7\t\t& 0.6\t\t&\t\t\t\t\t&\t\t&\t\t\t\t\t&\t\t&\t\\multicolumn{1}{l}{17.7}\t\t\t& $-$.05\t&\t\t\t\t\t\t&\t\t&\tR-IIId$^{e}$  \\\\ 18/04/07\t& 4208.7\t& 2.0\t\t& 1.9\t\t&\t\t\t\t\t&\t\t&\t\t\t\t\t&\t\t&\t\\multicolumn{1}{l}{17.8}\t\t\t& $-$.05\t&\t\t\t\t\t\t&\t\t&\tR-IIIb$^{e}$ \\\\ 23/04/07\t& 4214.1\t& 7.4\t\t& 7.2\t\t& 19.93 {\\tiny $\\pm$ .09}\t& $-$.18\t& 18.98 {\\tiny $\\pm$ .04}\t& $-$.14\t&\t18.38 {\\tiny $\\pm$ .05}\t\t\t& $-$.06\t&\t18.16 {\\tiny $\\pm$ .11}\t& $-$.01\t&\tLT \\\\ 27/04/07\t& 4218.4\t& 11.7\t& 11.3\t\t\t& 20.38 {\\tiny $\\pm$ .17}\t& $-$.18\t& 19.35 {\\tiny $\\pm$ .13}\t& $-$.15\t&\t18.95 {\\tiny $\\pm$ .09}\t\t\t& $-$.07\t&\t18.67 {\\tiny $\\pm$ .10}\t& $-$.01\t&\tFTN \\\\ 28/04/07\t& 4218.9\t& 12.2\t& 11.9\t\t\t& 20.48 {\\tiny $\\pm$ .16}\t& $-$.18\t& 19.52 {\\tiny $\\pm$ .07}\t& $-$.15\t&\t18.95 {\\tiny $\\pm$ .12}\t\t\t& $-$.07\t&\t18.69 {\\tiny $\\pm$ .07}\t& $-$.01\t&\tLT \\\\ 02/05/07\t& 4223.3\t& 16.7\t& 16.1\t\t\t& 21.03 {\\tiny $\\pm$ .41}\t& $-$.19\t& 19.93 {\\tiny $\\pm$ .11}\t& $-$.16\t&\t19.35 {\\tiny $\\pm$ .12}\t\t\t& $-$.08\t&\t19.02 {\\tiny $\\pm$ .09}\t& $-$.01\t&\tFTN \\\\ 05/05/07\t& 4225.9\t& 19.3\t& 18.6\t\t\t& 21.07 {\\tiny $\\pm$ .10}\t& $-$.19\t& 20.09 {\\tiny $\\pm$ .09}\t& $-$.15\t&\t19.38 {\\tiny $\\pm$ .07}\t\t\t& $-$.08\t&\t19.25 {\\tiny $\\pm$ .05}\t& $-$.01\t&\tLT \\\\ 11/05/07\t& 4232.6\t& 25.4\t& 24.5\t\t\t& 21.24 {\\tiny $\\pm$ .09}\t& $-$.22\t& 20.36 {\\tiny $\\pm$ .11}\t& $-$.12\t&\t19.69 {\\tiny $\\pm$ .08}\t\t\t& $-$.08\t&\t                        & \t\t&\tTNG \\\\ 13/05/07\t& 4234.0\t& 27.4\t& 26.4\t\t\t& 21.27 {\\tiny $\\pm$ .09}\t& $-$.22\t& 20.47 {\\tiny $\\pm$ .08}\t& $-$.12\t&\t19.72 {\\tiny $\\pm$ .06}\t\t\t& $-$.08\t&\t19.66 {\\tiny $\\pm$ .12}\t& $-$.01\t&\tLT \\\\ 19/05/07\t& 4240.0\t& 33.3\t& 32.2\t\t\t& 21.39 {\\tiny $\\pm$ .11}\t& $-$.21\t& 20.54 {\\tiny $\\pm$ .06}\t& $-$.11\t&\t20.02 {\\tiny $\\pm$ .07}\t\t\t& $-$.08\t&\t                  \t\t& \t\t&\tLT \\\\ 05/06/07\t& 4257.0\t& 50.3\t& 48.6\t\t\t& 21.46 {\\tiny $\\pm$ .06}\t& $-$.18\t& 20.61 {\\tiny $\\pm$ .20}\t& $-$.09\t&\t\t\t\t\t\t\t\t\t&\t\t&\t\t\t\t\t\t\t&\t\t&   LT \\\\ 13/06/07\t& 4265.0\t& 58.3\t& 56.4\t\t\t& 21.55 {\\tiny $\\pm$ .09}\t& $-$.17\t& 20.67 {\\tiny $\\pm$ .08}\t& $-$.07\t&\t20.32 {\\tiny $\\pm$ .10}\t\t\t& $-$.12\t&\t20.02 {\\tiny $\\pm$ .19}\t& .00\t&\tLT \\\\ 07/07/07\t& 4289.0\t& 82.3\t& 79.6\t\t\t& 21.67 {\\tiny $\\pm$ .07}\t& $-$.13\t& 20.93 {\\tiny $\\pm$ .08}\t& $-$.02\t&\t20.74 {\\tiny $\\pm$ .08}\t\t\t& $-$.15\t&\t20.47 {\\tiny $\\pm$ .14}\t& .00\t&\tLT \\\\ 12/08/07\t& 4324.9\t& 117.4\t& 113.4\t\t\t& 22.12 {\\tiny $\\pm$ .08}\t& $-$.07\t& 21.61 {\\tiny $\\pm$ .09}\t& .05\t&\t21.32 {\\tiny $\\pm$ .13}\t\t\t& $-$.20\t&\t20.95 {\\tiny $\\pm$ .13}\t& .00\t&\tWHT \\\\ \\hline \\end{tabular} \\end{center} {\\it Notes.} Errors on all magnitudes have been determined by taking into account the uncertainty of PSF fitting of the SN, the uncertainty of having accurately removed the background contamination measured via artificial-star experiments and the uncertainty of both calculating the colour-terms and zero-points of each night's imaging via the measurement of a local standard stars in the field of the SN. \\\\ $^{a}$ Relative to the date of discovery (CBET 927) \\\\ $^{b}$ {\\it K}-corrections determined for all epochs with spectral information and extrapolated to all other epochs (decimals). \\\\ $^{c}$ R-IIIb =  0.45m ROTSE-IIIb,  McDonald Observatory, Texas (USA). R-IIId =  0.45m ROTSE-IIId, Turkish National Observatory, Bakirlitepe (Turkey). LT = RATCam on the 2.0m Liverpool Telescope, Observatorio del Roque de Los Muchachos, La Palma (Spain). FTN = HawkCam on the 2.0m Faulkes Telescope North, Haleakala Observatory, Hawaii (USA). TNG = DOLoRes on the 3.58m Telescopio Nazionale Galileo, Observatorio del Roque de Los Muchachos, La Palma (Spain). WHT = auxiliary port camera on the 4.2m William Herschel Telescope, Observatorio del Roque de Los Muchachos, La Palma (Spain).\\\\ $^{d}$ Upper magnitude limit taken from CBET 927.\\\\ $^{e}$ Magnitude taken from CBET 927. \\end{minipage} \\end{table*}   \\begin{table*}\t\t%%%%%%%%%  TABLE OF PHOTOMETRY FOR SN 2007bi  %%%%%%%%%%%%%%%%%%%%%% \\setlength{\\tabcolsep}{5pt} \\begin{minipage}[t]{1.0\\textwidth} \\begin{center} \\caption{\\textrm {Photometry of SN 2007bi.}\\label{2007bi_phot}} \\begin{tabular}{l r r r l r l r l r l r l} \\hline\\hline {UT Date}\t& {JD}\t\t& Epoch$^{a}$\t& Rest-frame\t\t&{\\it B}\t& {$K_{BB}$}$^{b}$\t&{\\it V}\t& {$K_{VV}$}$^{b}$\t& {\\it R}\t& {$K_{RR}$}$^{b}$\t& {\\it I}\t& {$K_{II}$}$^{b}$\t& {Source}$^{c}$\t\\\\ & {- 2,450,000}\t& (d)\t\t\t& epoch (d)$^{a}$\t&\t\t&  \t\t\t\t&\t\t&\t\t\t\t&\t\t&\t\t\t\t&\t\t&\t\t\t\t&\t\t\t\t\\\\ \\hline 21/02/07 & 4153.0\t& 0.0\t\t& 0.0\t\t&\t\t\t\t\t\t\t\t&\t\t&\t\t\t\t\t\t\t\t&\t\t& \\multicolumn{1}{l}{17.64}\t& \t\t\t\t& \t\t\t\t\t\t&\t\t& G09$^{d}$ \\\\ 06/04/07 & 4197.0\t& 44.00\t\t& 39.04\t\t& \t\t\t\t\t\t\t\t&\t\t& \t\t\t\t\t\t\t\t&\t\t& \\multicolumn{1}{l}{18.3}\t\t& .00\t\t& \t\t\t\t\t\t&\t\t& POS$^{e}$ \\\\ 22/04/07 & 4213.6\t& 60.59\t\t& 53.76\t\t& 18.83 {\\tiny $\\pm$ .06}\t\t& $-$.36\t& 18.50 {\\tiny $\\pm$ .10}\t\t& $-$.09\t& 18.27 {\\tiny $\\pm$ .04}\t\t& $-$.03\t& 18.01 {\\tiny $\\pm$ .07}\t\t& $-$.20\t& Ekar \\\\ 23/04/07 & 4214.0\t& 61.02\t\t& 54.14\t\t& 18.86 {\\tiny $\\pm$ .03}\t\t& $-$.35\t& 18.45 {\\tiny $\\pm$ .03}\t\t& $-$.10\t& 18.34 {\\tiny $\\pm$ .02}\t\t& $-$.04\t& 18.03 {\\tiny $\\pm$ .05}\t\t& $-$.16\t& LT\\\\ 02/05/07 & 4223.3\t& 70.34\t\t& 62.41\t\t& 19.07 {\\tiny $\\pm$ .07}\t\t& $-$.32\t& 18.60 {\\tiny $\\pm$ .04}\t\t& $-$.12\t& 18.35 {\\tiny $\\pm$ .02}\t\t& $-$.08\t& 17.97 {\\tiny $\\pm$ .04}\t\t& $-$.10\t& FTN \\\\ 05/05/07 & 4226.0\t& 72.96\t\t& 64.74\t\t& 19.03 {\\tiny $\\pm$ .03}\t\t& $-$.32\t& 18.65 {\\tiny $\\pm$ .03}\t\t& $-$.12\t& 18.41 {\\tiny $\\pm$ .04}\t\t& $-$.08\t& 17.96 {\\tiny $\\pm$ .12}\t\t& $-$.12\t& LT \\\\ 12/05/07 & 4233.0\t& 79.97\t\t& 70.96\t\t& 19.20 {\\tiny $\\pm$ .02}\t\t& $-$.31\t& 18.69 {\\tiny $\\pm$ .02}\t\t& $-$.11\t& 18.51 {\\tiny $\\pm$ .02}\t\t& $-$.09\t& 18.10 {\\tiny $\\pm$ .05}\t\t& $-$.15\t& LT \\\\ 18/05/07 & 4239.3\t& 86.31\t\t& 76.58\t\t& 19.23 {\\tiny $\\pm$ .04}\t\t& $-$.30\t& 18.76 {\\tiny $\\pm$ .03}\t\t& $-$.11\t& 18.54 {\\tiny $\\pm$ .04}\t\t& $-$.10\t& 18.20 {\\tiny $\\pm$ .07}\t\t& $-$.19\t& FTN \\\\ 19/05/07 & 4234.0\t& 86.94\t\t& 77.14\t\t& 19.27 {\\tiny $\\pm$ .03}\t\t& $-$.30\t& 18.83 {\\tiny $\\pm$ .03}\t\t& $-$.11\t& 18.57 {\\tiny $\\pm$ .03}\t\t& $-$.10\t& \t\t       \t\t\t\t& \t\t& LT \\\\ 31/05/07 & 4251.9\t& 98.92\t\t& 87.77\t\t& 19.67 {\\tiny $\\pm$ .17}\t\t& $-$.29\t& 18.96 {\\tiny $\\pm$ .10}\t\t& $-$.10\t& 18.86 {\\tiny $\\pm$ .14}\t\t& $-$.11\t& 18.48 {\\tiny $\\pm$ .12}\t\t& $-$.25\t& LT \\\\ 07/07/07 & 4288.9\t& 135.92\t& 120.60\t& 20.43 {\\tiny $\\pm$ .02}\t\t& $-$.32\t& 19.68 {\\tiny $\\pm$ .03}\t\t& $-$.10\t&     \t\t\t\t\t\t&\t\t& 18.80 {\\tiny $\\pm$ .09}\t\t& $-$.15\t& LT \\\\ 12/08/07 & 4324.0\t& 171.88\t& 152.51\t& 20.73 {\\tiny $\\pm$ .07}\t\t& $-$.37\t& 19.94 {\\tiny $\\pm$ .12}\t\t& $-$.11\t& 19.64 {\\tiny $\\pm$ .08}\t\t& $-$.16\t& 19.06 {\\tiny $\\pm$ .11}\t\t& $-$.04\t& WHT \\\\ 10/04/08 & 4566.2\t& 413.19\t& 366.63\t& $>$22.59$^{f}$\t\t\t\t& $-$.54\t& 22.82 {\\tiny $\\pm$ .24}\t\t& $-$.18\t& 22.23 {\\tiny $\\pm$ .22}\t\t& $-$.07\t& 21.89 {\\tiny $\\pm$ .07}\t\t& $-$.08\t& VLT \\\\ 29/04/09 & 4951.1\t& 754.10\t& 669.12\t& $\\gtrsim$22.7\t\t\t\t\t&\t\t\t& $\\gtrsim$22.9\t\t\t\t\t&\t\t\t& $\\gtrsim$22.6\t\t\t\t\t\t\t\t\t\t\t\t&\t\t\t\t\t\t\t\t& $\\gtrsim$22.2\t\t\t\t\t\t\t\t\t\t\t\t&\t\t\t\t\t\t\t\t& LT \\\\ %10/04/08 & 4566.2\t& 369.2\t\t& 327.6\t\t& \t\t\t\t\t\t\t\t& $-$54\t& 22.32 {\\tiny $\\pm$ .05}\t\t& $-$18\t& 22.11 {\\tiny $\\pm$ .07}\t\t& $-$07\t& 21.89 {\\tiny $\\pm$ .07}\t\t& $-$08\t& VLT \\\\ %02/07/08 & 4649.9\t& 452.9\t\t& 401.9\t\t& 23.56 {\\tiny $\\pm$ .13}\t\t& $-$60\t& 22.93 {\\tiny $\\pm$ .16}\t\t& $-$20\t& 22.52 {\\tiny $\\pm$ .11}\t\t& $-$03\t& \t\t\t\t\t\t& \t\t& WHT \\\\ \\hline \\end{tabular} \\end{center} {\\it Notes.} Errors on all magnitudes have been determined by taking into account the uncertainty of PSF fitting of the SN, the uncertainty of having accurately removed the background contamination measured via artificial-star experiments and the uncertainty of both calculating the colour-terms and zero-points of each night's imaging via the measurement of a local standard stars in the field of the SN. \\\\ $^{a}$ Relative to the date of maximum brightness \\citep{Gal-Yam:2009}. \\\\ $^{b}$ {\\it K}-corrections determined for all epochs with spectral information and extrapolated to all other epochs (decimals). \\\\ $^{c}$ POS = QUEST II camera on the Palomar Oschin 1.2m Schmidt telescope, Palomar Observatory, Califonia (USA).  Ekar = AFSOC on the 1.82m Copernico Telescope, INAF, Osservatorio di Asiago, Mt. Ekar, Asiago (Italy). LT = RATCam on the 2.0m Liverpool Telescope, Observatorio del Roque de Los Muchachos, La Palma (Spain). FTN = HawkCam on the 2.0m Faulkes Telescope North, Haleakala Observatory, Hawaii (USA). TNG = DOLoRes on the 3.58m Telescopio Nazionale Galileo, Observatorio del Roque de Los Muchachos, La Palma (Spain). WHT = auxiliary port camera on the 4.2m William Herschel Telescope, Observatorio del Roque de Los Muchachos, La Palma (Spain). VLT = FORS2 on the 8.2m Very Large Telescope (unit telescope 1), Paranal Observatory, Cerro Paranal in the Atacama Desert (Chile).\\\\ $^{d}$ Peak magnitude determined by \\cite{Gal-Yam:2009} via parabolic fitting to pre- and post-peak magnitude data. $^{e}$ Magnitude taken from CBET 929. $^{f}$ Upper-limit: flux dominated by host-galaxy. \\end{minipage} \\end{table*} ",
        "conclusions": "We have presented {\\it BVRI} photometric and optical spectroscopic data for SN 2007bg and SN 2007bi. Neither SNe 2007bg nor 2007bi were found in association with a GRB, but from estimates of the metallicities of their host-galaxies they are found to inhabit similar low-metallicity environments to GRB/SNe Ic. SN 2007bg is found to be hosted by an extremely sub-luminous galaxy of magnitude $M_{B}=-12.44${\\tiny$\\pm0.78$} mag with a measured oxygen abundance of 12+log(O/H$)_{\\rm PP04}=8.18${\\tiny $\\pm0.17$}. SN 2007bg also displays one of the fastest post-maximum decline rates of all SNe BL-Ic. This fast light curve decline combined with high expansion velocities gives SN 2007bg a high kinetic energy to ejected mass ratio ($E_K/M_{ej}\\sim2.7$).  Reaching a peak magnitude of M$_{R}\\sim-21.3$ mag, the light curve of SN 2007bi displays a remarkably slow decline, following the decay rate of $^{56}$Co to $^{56}$Fe throughout the course of its observed lifetime. SN 2007bi also displays an extreme longevity in its spectral evolution and is still not fully nebular at approximately one year post-maximum. From modeling the bolometric light curve of SN 2007bi we estimate a total ejected $^{56}$Ni mass of $M_{Ni}=3.5-4.5{\\rm M}_\\odot$, the largest $^{56}$Ni mass measured in the ejecta of a SN to date. The spectra are characterised by the presence of strong oxygen and iron lines, suggesting the explosion of a massive stellar core which synthesised a remarkable amount of $^{56}$Ni. A strong case can be made that SN 2007bi was a pair-instability supernova, but given that it is hosted by only a moderately low-metallicity galaxy we suggest that a less-exotic core-collapse mechanism may have been the cause of this explosion.  In the future we hope to obtain a spectrum of the host-galaxy of SN 2007bi, now that the SN has faded, in order to more accurately measure the oxygen abundance. It will also be of value to reobserve the host of SN 2007bg to get a measurement of the oxygen lines required to determine a $T_e$ oxygen abundance. It may be that the low-metallicities of the host-galaxies of these two SNe Ic have been one of the main parameters contributing to their peculiarity. With the advent of current and future all-sky surveys such as Pan-STARRS, Skymapper, PTF and RSVP we are certain to find many more of these strange and interesting objects hosted in low-metallicity environments \\citep[see][]{Young:2008p8179}."
    },
    "0910/0910.3615_arXiv.txt": {
        "abstract": "{ We investigate supersymmetric models for dark matter which is represented by pseudomoduli in weakly coupled hidden sectors.  We propose a scheme to add a dark matter sector to quiver gauge theories with metastable supersymmetry breaking.  We discuss the embedding of such scheme in string theory and we describe the dark matter sector in terms of D7 flavour branes.  We explore the phenomenology in various regions of the parameters.}       \\begin{document} ",
        "introduction": "Cosmological observations have established the existence of dark matter which is not composed by any of the Standard Model (SM) particles, and with a relic abundance of the order of $\\Omega h^2 \\sim 0.1$ (for reviews see \\cite{Jungman:1995df,Trodden:2004st,Rubakov:2005tx,Hooper:2009zm} and reference therein).  A stable particle in thermal equilibrium with the SM in the early universe is usually referred as cold dark matter (DM). As long as the universe expands, the DM ceases to annihilate efficiently and freezes out, leaving a relic abundance \\cite{relicosi} \\be \\Omega h^2 \\sim \\frac{1}{\\langle  \\sigma v \\rangle} \\sim \\frac{\\alpha_{DM}}{M_{DM}^2} \\ee where $h$ is the Hubble parameter in units of 100 km/sec per Mpc, $\\langle \\sigma v \\rangle$ is the annihilation cross section, $\\alpha_{DM}$ is the characteristic coupling and $M_{DM}$ is the mass of the dark matter particle. For a massive particle with weak interaction ($g_{DM} \\sim 1$), i.e. a WIMP, the required relic abundance is obtained for $M_{DM} \\sim 1$ TeV. The TeV scale that naturally appears suggests some relations between DM and electroweak symmetry breaking (EWSB), and has been dubbed as the WIMP miracle.  A standard explanation for the origin of the electroweak scale is supersymmetry breaking (for a phenomenological introduction to supersymmetry and references see \\cite{Martin:1997ns}). The idea that supersymmetry and its breaking should account also for the origin of the DM has prompted an extensive investigation. Much effort has been focused on the mechanism of gravity mediation, where the lightest neutralino of the MSSM should be a stable, viable DM candidate. A phenomenological drawback for models of gravity mediation is that they do not address the susy flavour problem.  The flavour problem is solved in the gauge mediation scenario \\cite{Dine:1982zb,Giudice:1998bp}, because the SM gauge interactions are flavour blind. In gauge mediated models supersymmetry is broken at low energies and the gravitino is the LSP. However, the gravitino is typically too light (with mass in the range of eV to Gev) \\cite{Fayet:1977vd,Pagels:1981ke} to be a viable cold DM candidate. Other alternatives for susy DM have been and are under inspection (see for instance \\cite{sviluppi,Pospelov:2007mp,Chun:2008by}).  In gauge mediated theories particles in the hidden sectors can be alternative DM candidates. In such models, supersymmetry breaking in the hidden sector occurs dynamically, via strong dynamics effects. The DM can then be identified with the mesons and the baryons of the strongly coupled hidden sector \\cite{Dimopoulos:1996gy}.  On the other hand, one can explore the possibility of realizing dark matter in weakly coupled hidden sector with spontaneous supersymmetry breaking \\cite{Shih:2009he,KerenZur:2009cv}. These models are relevant since they can arise as the low energy description of UV free gauge theories, as shown by \\cite{Intriligator:2006dd}.  Such weakly coupled models typically break supersymmetry in metastable vacua, with a spectrum of elementary particles. The scalar potential often presents tree level flat directions, which are lifted by one loop quantum corrections. These pseudomoduli fields have weak scale interactions and their one loop mass is of the order of the supersymmetry breaking scale, i.e the TeV scale naturally arises. If they are stable against decay because of a discrete $Z_2$ symmetry, they represent viable cold dark matter candidates. This possibility has been explored in O'Raifeartaigh like models in \\cite{Shih:2009he} and also analyzed in \\cite{KerenZur:2009cv}.  In this paper we shall study models of pseudomoduli as dark matter and wish to answer some questions concerning the relative UV completion. This is a non trivial problem, with severe constraints \\cite{Adams:2006sv,Antoniadis:2007xc} for consistent embedding in a UV complete theory.  We shall propose a quiver model for which a stringy origin can be reached.  We embed models with pseudomoduli DM in appropriate quiver gauge theories which arise from $D$-branes at CY singularities.  The starting point is a quiver theory, inherited and/or interpreted as an IR Seiberg dual, which provides the hidden sector of dynamical breaking of susy.  The next step is the addition of a dark matter sector, represented by an extra node in the quiver connected through matter interactions to the hidden sector.  The addition of $D7$ flavour branes in the CY can be put in correspondence with the dark matter sector in the quiver.   The structure of the paper is the following. In section \\ref{pseudomoduliDM} we propose the general strategy to embed pseudomoduli dark matter in quiver gauge theories, we comment on the phenomenological constraints and also on the string theory realization.  In section \\ref{KOOsec} we build a concrete example; we couple pseudomoduli DM to the KOO model \\cite{Kitano:2006xg}, and we discuss the related phenomenology.  In section \\ref{stringemb} we provide for a UV origin to the model in terms of a step of Seiberg duality: the UV model is obtained by deforming an $L^{131}$ non isolated singularity and by adding flavour D7 branes.  A conclusion follows. In appendix \\ref{decay} we review a basic cosmological bound on the supersymmetry breaking scale.  In appendix \\ref{albeapp} we review the procedure of flavoring with D7 branes and we discuss its relation with the DM sector.  In appendix \\ref{APP} we provide the details of the one loop computations for the pseudomoduli masses.\\\\ ~\\\\ ~\\\\ As we were finishing this paper, we were informed of \\cite{Inaki} which study leptophilic dark matter \\cite{Fox:2008kb} in quiver gauge theories. ",
        "conclusions": "In this article we proposed a scheme to realize pseudomoduli dark matter in quiver gauge theories with weakly coupled supersymmetry breaking a' la ISS. In section \\ref{pseudomoduliDM} we distinguished a metastable supersymmetry breaking sector and a dark matter sector. The former communicates the breaking of supersymmetry to the MSSM trough gauge interactions. This mechanism requires $R$-symmetry to be broken, spontaneously or explicitly. The dark matter sector contains two pseudomoduli fields and it is coupled to the supersymmetry breaking sector through superpotential terms. This coupling induces, at one loop, supersymmetry breaking masses for the pseudomoduli and their fermionic partners in the dark matter sector. This provides the mechanism to generate a TeV scale mass for the fermions which are the cold dark matter candidates.  We considered only the case of uncharged dark matter with respect to the gauge interactions of the standard model. There is a coupling of the $U(1)_{d}$ gauge group, under which DM is charged, to the $U(1)_Y$ of the standard model via the kinetic mixing. Cases in which the dark matter is charged under a non abelian gauge group are left for future works.   We showed that when the supersymmetry breaking sector is described by $D3$ branes probing CY singularity, the dark matter sector corresponds to the addition of $D7$ flavour branes.  In section \\ref{KOOsec} we realized the proposed scenario by coupling a pseudomoduli DM sector to the KOO model, and $1$TeV mass dark matter is rather natural. In section \\ref{stringemb} we studied a model with D$3$ and D$7$ branes wrapped over a deformed CY $L^{131}$ singularity. We have argued that this model reduces, in the infrared, to the KOO model plus DM.  Many extensions of our proposal can be studied. One can find other models by coupling the DM sector to gauge theories different from KOO. For example one can couple DM to models with  $R$-symmetry breaking in metastable vacua \\cite{Amariti:2006vk, Csaki:2006wi,Abel:2007jx,Intriligator:2007py, Shih:2007av,Giveon:2007ef,Haba:2007rj, Giveon:2008ne,Essig:2008kz}. Another interesting issue is the investigation at large of the UV completion. The generation of the DM sector inside a quiver gauge theory can be a result of a step of Seiberg duality, as we showed in the deformed and flavored $L^{131}$ theory."
    },
    "0910/0910.5269.txt": {
        "abstract": "We present  polarization observations of the dust continuum emission from the young star forming region IRAS~16293. These observations of IRAS~16293, which is a binary system, were conducted by the Submillimeter Array (SMA) at an observing frequency of 341.5 GHz ($\\lambda \\sim 880\\mu$m) and with high angular resolution ($\\sim$2$\\arcsec$--3$\\arcsec$).  We find that the large scale global direction of the field, which is perpendicular to the observed polarization, appears to be along the dust ridge where the emission peaks.   On smaller scales we find that the field structure is significantly different for the two components of the binary. The first component, source A, shows a magnetic field structure which is ``hourglass'' shaped as predicted from theoretical models of low mass star formation in the presence of strong magnetic fields. However, the other component, source B, shows a relatively ordered magnetic field with no evidence of any deformation.  We have possibly detected a third younger outflow from source A as seen in the SiO emission which is in addition to the two well known powerful bipolar outflows in this kinematically active region.  There is an observed decrease in polarization towards the center and this ``polarization hole'' is similar to decreases seen in other young star forming regions.  Our calculations show that in IRAS~16293 the magnetic energy is stronger than the turbulent energy but is approximately similar to the centrifugal energy. There is considerable misalignment between the outflow direction and the magnetic field axis and this is roughly in agreement with model predictions where the magnetic energy is comparable to the centrifugal energy. In conjunction with other observations of the kinematics  as determined from the outflow energetics and chemical differentiation we find that our results provide additional evidence to show that the two protostars appear to be in different stages during their evolution. ",
        "introduction": "In the ``classical picture'' of star formation, magnetic fields are believed to strongly influence star formation activity in molecular clouds. They provide support to a cloud against gravitational collapse and thus explain the low efficiency of the star formation process \\citep{Mouschovias01, Shu07}. The process of ambipolar diffusion, in which magnetic flux is redistributed in the cloud, leads to the formation of a core that can no longer be magnetically supported and gravitational collapse sets in. In addition, the process of magnetic braking can help to remove angular momentum and slow down the rotation of the cloud as it collapses \\citep{Basu94}. In contrast, there have been a number of alternate theories which postulate that magnetic fields are relatively weak and supersonic magnetohydrodynamic turbulence is the dominant process \\citep{MacLow04}. Turbulence controls the evolutions of clouds, and cores form at the intersection of supersonic turbulent flows. Only a fraction of such cores become supercritical and collapse begins to occur on a gravitational free-fall timescale \\citep{Ballesteros07}.  Studying the morphology of the magnetic field in the interstellar medium (ISM) requires difficult and sensitive observations of polarized radiation \\citep{Hildebrand04}.  Spinning dust grains in the ISM become partially aligned with the magnetic field, generally with their long axes perpendicular to the field \\citep{Davis51}. Nevertheless, the exact nature of the alignment process is a matter of debate \\citep{Lazarian07}, but it appears that the dominant mechanism may actually be alignment by radiative torques \\citep{Hoang08, Hoang09}.  As a consequence of this alignment, the thermal dust emission is partially linearly polarized, with polarization direction perpendicular to the magnetic field.  The magnitude of the field can be determined indirectly by using the dispersion of the position angles of the field under the assumption of energy equipartition between the kinetic and perturbed magnetic energies \\citep{Chandra53, Ostriker01,Heitsch01,Crutcher04,Falceta08}.  Observations of the large scale polarization distribution in molecular clouds  from optical, infra-red, and submillimeter (submm) wavelengths usually show an ordered pattern \\citep{Dotson00, Pereyra04, Poidevin06, Matthews09}. The analysis of polarization data has showed that some dark clouds are  magnetically dominated \\citep{Alves08, Heyer08} while other giant molecular clouds are close to equipartition with turbulence \\citep{Novak09}.  However, observationally, the debate on whether turbulence or magnetic fields are dominant is still unresolved \\citep{Crutcher09}.  Early and pioneering work with the BIMA millimeter~(mm) array showed the dust (and molecular line) polarized emission can be traced at an angular resolution of few arcseconds, the scales where the core contraction process occurs \\citep{Rao98,Girart99,Lai01,Lai03,Cortes05}. Currently, the Submillimeter Array (SMA) is the only telescope that can detected the polarized emission at mm and submm wavelengths  \\citep{Marrone06b, Girart06, Girart09,Tang09a, Tang09b}. In the object NGC~1333~IRAS~4A (IRAS~4A here after), the SMA observations have revealed that  the magnetic field configuration is consistent with theoretical models for the formation of solar-type stars in which magnetic fields play a much stronger role than turbulence \\citep{Girart06, Goncalves08}.  The low-mass star forming region IRAS~$16293-2422$ (IRAS~16293 here after), located in the $\\rho$ Ophiuchi molecular cloud complex, has been the focus of numerous studies since the report of infall spectral signatures \\citep{Walker86,Menten87}.  Accurate observations using astrometric techniques show that the $\\rho$ Ophiuchi molecular cloud is at a distance $\\simeq 120$~pc \\citep{Knude98,Loinard08}, while maser VLBI observations towards IRAS~16293 suggest a distance of $\\simeq 178$~pc \\citep{Imai07}.  For this paper we have assumed a distance of $d=150$~pc.  IRAS~16293 is a Class~0 protostellar system with a bolometric luminosity of 32~L$_{\\sun}$, and is surrounded by a compact, $R_{\\rm env} \\simeq3000$~AU, but relatively massive envelope, $\\simeq3.0$~M$_{\\sun}$, \\citep{Correia04}. Higher resolution radio interferometric observations have revealed a double core separated by  5$\\arcsec$ (750~AU in the plane of the sky). The southern core is commonly referred to as source A and the northern one as source B. Both cores appear to have different physical and chemical properties \\citep{Wootten89, Estalella91, Bottinelli04,Chandler05,Takakuwa07}.  Despite its low luminosity, IRAS~16293 has a rich chemistry, with hot-core like properties at scales of $\\sim100$~AU  \\citep{Blake94, Ceccarelli00, Schoier02, Cazaux03, Kuan04, Bisschop08}. There is a strong quadrupolar outflow associated with source~A \\citep{Walker88,Mizuno90,Stark04,Yeh08}. The larger bipolar outflow is in the east-west direction (E-W outflow), with a position angle (PA) of $110\\arcdeg$. The other outflow is in the northeast-southwest direction (NE-SW outflow), with a PA of $60\\arcdeg$ and shows copious SiO emission \\citep{Hirano01}. High resolution submm observations \\citep{Chandler05} have revealed multiplicity in source A and which of these sources powering these two outflows is a matter of debate \\citep{Loinard07}.  Previous measurements of the magnetic field geometry in this object by observing polarized dust emission have yielded contradictory results \\citep{Flett91,Tamura93,Akeson97}, possibly because they were sensitivity limited.  Thus, it is apparent that more sensitive and higher resolution observations of the magnetic field structure are needed.  In \\S~2 we briefly describe the observations and the data reduction procedure. \\S~3 presents the results, \\S~4 the analysis, and \\S~5 discusses the possible scenarios for the observed magnetic field morphology in context with the information from observations already known from the literature. ",
        "conclusions": "The installation of a polarimetry system on the SMA has enabled us to map the magnetic field structure in the ISM, especially in young star forming regions.  Using this system, we have obtained high angular resolution and high sensitivity maps of the magnetic field structure in IRAS~16293 through observations of the polarized dust continuum. These observations significantly improve on the earlier measurements which detected extremely low fractional polarization. Our detections indicate that those early attempts were limited in their sensitivity and angular resolution.  At the same time, the net polarization over the entire region mapped by us is quite small ($\\sim$ 0.2\\%) and this is indeed in agreement with past measurements. It is also apparent that the polarization fraction appears to decrease towards the center where the intensity of the source increases. However, it seems to taper off at sufficiently large values of the intensity especially in source B. This is possibly due to the effects of limited resolution towards the center where the peak of the emission occurs. Further higher resolution observations will be needed to determine if this relationship continues even at shorter distances from the center.  The two sources A and B have significantly different magnetic field morphologies.  In source A, our maps show that the magnetic field has the pinched ``hourglass'' shape that is expected from theoretical calculations.  In contrast, the field lines in source B, appear to be quite uniform. Using the Chandrasekhar-Fermi method to obtain the  field strength and the continuum dust emission to calculate the mass, we calculate the mass-to-flux ratio to be $\\sim 1$.  This implies that this object is in (or close to) the supercritical stage and the magnetic field is no longer able to prevent the collapse. In addition, the ordered field structures indicate that the magnetic energy likely exceeds the turbulent energy, but is comparable to the centrifugal energy. However, there are other reported measurements, that are not as sensitive, which appear to show  that turbulence increases on more intermediate sizescales within the cloud. Nevertheless, the largest scales, which are probed in absorption polarimetry, once again show fairly ordered magnetic field structures.  The evolutionary stages and the ages of the two sources are likely to be different and the SMA observations of the magnetic field structure also appear to confirm this view.  The physical and chemical properties also differ in the two sources.  In source B, most of the mass at scales of a few hundred AU has already been accreted onto a compact, optically thick disk, while source A still has a significant fraction of its mass in it's circumstellar envelope. Furthermore, the nature of the outflow activity is quite different in the two sources: source A shows significant activity and drives at least two powerful outflows and possibly a third compact outflow which we have detected in SiO emission, while source B shows no ongoing activity except for a possible fossil or remnant outflow. Furthermore, source A appears to have a richer chemical environment than source B.  It is possible that source B may not be a true Class~0 protostar and could possibly be a transitional object between Class~0 and Class~I.  There appears to be a strong misalignment between the outflow direction in source A and the magnetic field axis. This is in approximate agreement with theoretical model predictions when the magnetic energy is not significantly greater than the centrifugal energy.  This is the second such low-mass star forming region in which it has been shown that magnetic fields appear to dominate over turbulence. However, the magnetic field morphology has been mapped in only a small number of such star forming regions. Improvements in telescope sensitivity to the polarized flux density will allow us to significantly expand this sample. With the increase in the number of objects studied, it will be possible to get a clearer and statistically significant picture and thereby address the all important question: Which process plays a dominant role in the star forming process --- turbulence or magnetic fields?"
    },
    "0910/0910.3565_arXiv.txt": {
        "abstract": "We report the $Fermi$-LAT discovery of high-energy (MeV/GeV) $\\gamma$-ray emission positionally consistent with the center of the radio galaxy M87, at a source significance of over $10\\sigma$ in ten-months of all-sky survey data. Following the detections of Cen~A and Per~A, this makes M87 the third radio galaxy seen with the LAT. The faint point-like $\\gamma$-ray source has a $>$100 MeV flux of 2.45 ($\\pm 0.63$) \\gunit\\ (photon index = $2.26 \\pm 0.13$) with no significant variability detected within the LAT observation. This flux is comparable with the previous EGRET upper limit ($< 2.18$ \\gunit, 2$\\sigma$), thus there is no evidence for a significant MeV/GeV flare on decade timescales. Contemporaneous $Chandra$ and VLBA data indicate low activity in the unresolved X-ray and radio core relative to previous observations, suggesting M87 is in a quiescent overall level over the first year of $Fermi$-LAT observations. The LAT $\\gamma$-ray spectrum is modeled as synchrotron self-Compton (SSC) emission from the electron population producing the radio-to-X-ray emission in the core. The resultant SSC spectrum extrapolates smoothly from the LAT band to the historical-minimum TeV emission. Alternative models for the core and possible contributions from the kiloparsec-scale jet in M87 are considered, and can not be excluded. ",
        "introduction": "As one of the nearest radio galaxies to us \\citep[$D=16$ Mpc is adopted;][]{ton91}, M87 is amongst the best-studied of its source class. It is perhaps best known for its exceptionally bright arcsecond-scale jet \\citep{cur18}, well-imaged at radio through X-ray frequencies at increasingly improved sensitivity and resolution over the decades \\citep[e.g.,][]{bir91,spa96,mar02,per05}. Near its central $\\sim (3-6) \\times 10^{9}$ solar mass supermassive black hole \\citep{mac97,geb09}, the jet base has been imaged down to $\\sim$0.01 pc resolution \\citep[$\\sim 15-30\\times$ the Schwarzschild radius,][]{jun99,ly07,tev09}.  At the highest energies, M87 is regularly detected by HESS, MAGIC, and VERITAS with variable TeV emission on timescales of years and flaring in a few days \\citep{aha06,alb08,acc08,tev09}. The sensitivity of these Cherenkov telescopes have also enabled the detection of another well-known nearby radio galaxy Cen~A \\citep{aha09}. Without comparable imaging resolution to the lower energy studies however, variability and spectral modeling are necessary to infer the production site of the TeV $\\gamma$-rays and to deduce the source physical parameters.  At high energy $\\gamma$-rays ($\\sim$20 MeV -- 100 GeV), we are similarly poised for new radio galaxy discoveries with the Large Area Telescope (LAT) aboard the recently launched {\\it Fermi Gamma-ray Space Telescope} \\citep{atw09}. Indeed, we report here the detection of a faint, point-like $\\gamma$-ray source positionally coincident with M87 using the $Fermi$-LAT. After the confirmation of the EGRET discovery of Cen~A \\citep[][and in preparation]{sre99,lbas}, and the recent detection of Per~A/NGC~1275 \\citep{pera}, this is the third radio galaxy successfully detected by $Fermi$. Unlike the known variable TeV source, there is so far no evidence for variability of the MeV/GeV emission in M87. An origin of the LAT emission from the unresolved parsec scale jet (hereafter, denoted as the `nucleus' or `core') observed contemporaneously with $Chandra$ and the VLBA\\footnote{The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.} is discussed. Potential contributions from the larger-scale ($\\simgt 0.1-1$ kpc) jet to the unresolved $\\gamma$-ray source are also briefly considered. Section 2 contains the details of the LAT observations, including a description of the $Chandra$ and VLBA data utilized, with the discussion of these results in section 3. ",
        "conclusions": "\\label{section-discussion}  In blazars, it is commonly believed that $\\gamma$-rays are produced in compact emission regions moving with relativistic bulk velocities in or near the parsec scale core in order to explain the observed rapid variability and to avoid catastrophic pair-production \\citep[e.g.,][]{don95}. Consequently, it is natural to extend this supposition to radio galaxies \\citep{chi01} which are believed to have jets oriented at systematically larger angles to our line of sight, thus constituting the parent population of blazars. Indeed, in the case of M87, a significant months timescale rise in the flux of the sub-pc scale radio core was discovered with the VLBA (at 43 GHz) during a period in early 2008 when few day timescale TeV flaring was detected \\citep{tev09}. During this period of increased activity, a $Chandra$ measurement of the sub-arcsecond scale X-ray nucleus also indicated a relatively higher flux than seen in past observations \\citep{har09}, thus signaling a common origin for the flaring emissions in the M87 nucleus. Therefore, during periods of lower $\\gamma$-ray activity, the radio/X-ray core can also be considered a dominant source of the unresolved higher-energy emission, and we discuss this in the context of the LAT MeV/GeV detection.  In Figure~\\ref{figure-sedgamma}, the LAT spectrum of M87 is plotted along with representative integrated TeV spectra from HESS \\citep{aha06}. The TeV measurements cover periods when M87 was in its historical-minimum (in 2004), and during a high state \\citep[in 2005 -- cf., Fig.~3 in][]{acc08}. Although the formal difference in the fitted photon indices of the TeV data at high and low states is not statistically significant ($\\Gamma = 2.22 \\pm 0.15$ and $2.62 \\pm 0.35$, respectively), the LAT MeV/GeV spectrum ($\\Gamma =2.3$) connects smoothly with the low-state TeV spectrum. Taken together with the X-ray and radio measurements obtained during the LAT observation (\\S~\\ref{section-lat}), we view this as an indication that M87 is in an overall low $\\gamma$-ray activity state during the considered period. In fact, no significant TeV flaring was detected in a preliminary analysis of 18 hrs of contemporaneous VERITAS observations from Jan.\\ - Apr.\\ 2009 \\citep{hui09}.  M87 is the faintest $\\gamma$-ray radio galaxy detected so far by the LAT with a $>$100 MeV flux ($\\sim 2.5$\\gunit) about an order of magnitude lower than in Cen~A \\citep{lbas} and Per~A \\citep{pera}; the corresponding $>100$ MeV luminosity, $4.9 \\times 10^{41}$ erg s$^{-1}$, is 4$\\times$ greater than that of Cen~A, but $>$200$\\times$ smaller than in Per A. There is no evidence of intra-year or decade-timescale MeV/GeV variability in M87 (\\S~\\ref{section-lat}), in contrast to the $\\simgt 7\\times$ and $\\sim 1.6\\times$ larger observed LAT fluxes than the previous EGRET ones in the cases of Per~A \\citep{pera} and Cen~A \\citep{lbas}, respectively. The $\\gamma$-ray photon index of M87 in the LAT band is similar to that of Per~A ($\\Gamma = 2.3$ and 2.2, respectively), while being smaller than observed in Cen~A \\citep[$\\Gamma=2.9$;][]{lbas}. These sources are low-power (FRI) radio galaxies, and have broad low-energy synchrotron and high-energy inverse Compton (IC) components in their spectral energy distributions (SEDs) peaking roughly in the infrared and $\\gamma$-ray bands, respectively. Low-energy peaked BL Lac objects have similar shaped SEDs, with approximately equal apparent luminosities \\citep[e.g.,][]{kub98}. As FRI radio galaxies are believed to constitute the parent population of BL Lacs in unified schemes \\citep{urr95}, the overall similarity of their SEDs is not surprising.  We construct a SED for M87 (Figure~\\ref{figure-sed}) using the the LAT $\\gamma$-ray spectrum and the overlapping Jan.\\ 7, 2009 $Chandra$ and VLBA measurements of the core. Also plotted are historical radio to X-ray fluxes of the core \\citep[see][]{spa96,tan08} measured at the highest resolutions at the respective frequencies. The core is known to be variable, with factors of $\\sim2$ changes on months timescales common in the optical and X-ray bands \\citep{per03,har09}. To help constrain the overall SED at frequencies between the X-ray and LAT measurements, we determined integrated 3$\\sigma$ upper limits in three hard X-ray bands \\citep[following,][]{aje08} based on the $Swift$/BAT dataset in \\citet{aje09}, including about another additional year of exposure (i.e., $\\sim$4 years total from Mar.\\ 2005 - Jan.\\ 2009).   \\begin{figure} \\epsscale{1.2} \\plotone{fig4.eps} \\caption{SED of M87 with the LAT spectrum and the Jan.\\ 7, 2009 MOJAVE VLBA 15 GHz and $Chandra$ X-ray measurements of the core indicated in red. The non-simultaneous 2004 TeV spectrum described in Figure~\\ref{figure-sedgamma} and $Swift$/BAT hard X-ray limits (\\S~\\ref{section-discussion}) of the integrated emission are shown in light brown. Historical measurements of the core from VLA 1.5, 5, 15 GHz \\citep{bir91}, IRAM 89 GHz \\citep{des96}, SMA 230 GHz \\citep{tan08}, $Spitzer$ 70, 24 $\\mu$m \\citep{shi07}, Gemini 10.8 $\\mu$m \\citep{per01}, $HST$ optical/UV \\citep{spa96}, and $Chandra$ 1 keV from \\citet[][hidden behind the new measurements]{mar02} are plotted as black circles. The VLBA 15 GHz flux is systematically lower than the historical arcsec-resolution radio to infrared measurements due to the presence of intermediate scale emission \\citep[see e.g.,][]{kov07}. The blue line shows the one-zone SSC model fit for the core described in \\S~\\ref{section-discussion}. } \\label{figure-sed} \\end{figure}   The broad-band SED is fit with a homogeneous one-zone synchrotron self-Compton (SSC) jet model \\citep{fin08} assuming an angle to the line of sight, $\\theta$=10\\deg, and bulk Lorentz factor, $\\Gamma_{\\rm b} = 2.3$ (Doppler factor, $\\delta = 3.9$), consistent with observations of apparent motions of $\\simgt 0.4c$ ($\\Gamma_{\\rm b} > 1.1$) in the parsec-scale radio jet \\citep{ly07}. A broken power-law electron energy distribution $N(\\gamma)\\propto\\gamma^{-p}$ is assumed, and the indices, $p_{\\rm 1}=1.6$ for $\\gamma$ = [1, $4 \\times 10^{3}$] and $p_{\\rm 2}=3.6$ for $\\gamma$ = [$4 \\times 10^{3}$, $10^{7}$] are best guesses based on the available core measurements. The normalization at low energies is constrained by the single contemporaneous VLBA 15 GHz flux which is measured with $\\sim 10^{2}-10^{3}\\times$ better resolution than the adjacent points. The source radius, $r= 1.4 \\times 10^{16}$ cm = 4.5 mpc is chosen to be consistent with the best VLBA 43 GHz map resolution \\citep[$r<$7.8 mpc = 0.1 mas,][]{jun99,ly07} and is of order the size implied by the few day timescale TeV variability \\citep{tev09}. For the source size adopted, internal $\\gamma-\\gamma$ absorption is avoided so that the LAT spectrum extends relatively smoothly into the TeV band, consistent with the historical-minimum flux detected by HESS \\citep{aha06} and the preliminary upper limit of $<1.9\\%$ Crab from VERITAS observations \\citep{hui09} contemporaneous with the LAT ones.  In the SSC model, the magnetic field is $B = 55$ mG and assuming the proton energy density is 10$\\times$ greater than the electron energy density, the total jet power is $P_{\\rm j} \\sim 7.0 \\times 10^{43}$ erg s$^{-1}$. The jet power is particle dominated, with only a small contribution from the magnetic field component ($P_{\\rm B} \\sim 2 \\times 10^{40}$ erg s$^{-1}$). In comparison, the total kinetic power in the jet is $\\sim$few $\\times 10^{44}$ erg s$^{-1}$ as determined from the energetics of the kpc-scale jet and lobes \\citep{bic96}, and is consistent with the jet power available from accretion, $P_{\\rm j} \\simlt 10^{45}$ erg s$^{-1}$ \\citep{rey96,dim03}. These power estimates are similar to those derived for BL Lacs from similarly modeling their broad-band SEDs \\citep[e.g.,][]{cel08}.  As applied to M87, such single-zone SSC emission models also reproduce well the broad-band SEDs up to MeV/GeV energies in the radio galaxies Per~A \\citep{pera} and Cen~A \\citep{chi01}. In this context, the observed MeV/GeV $\\gamma$-ray fluxes of blazars appear to be correlated with their compact radio cores \\citep{lbas,kov09}, suggesting a common origin in the Doppler boosted emission in the sub-parsec scale jets. The fact that the 3 radio galaxies detected by the LAT so far have amongst the brightest ($\\simgt$1 Jy) unresolved radio cores, in line with these expectations \\citep{ghi05}, lend evidence for a common connection between the $\\gamma$-ray and radio emitting zones in such jets.  It should be emphasized that these observations are not simultaneous and particularly, the TeV emission is known to be variable on year timescales, so other emission components may contribute to the variable emission. Therefore, although not strictly required, more sophisticated models over the single-zone one presented can reproduce or contribute to the observed emission. In particular, the beaming requirements in the one-zone SSC modeling of the three known $\\gamma$-ray radio galaxies are systematically lower than required in BL Lacs, suggesting velocity profiles in the flow \\citep{chi01}. Such models \\citep{geo05,tav08} have in fact been used to fit the SED of M87 in addition to models based on additional spatial structure \\citep[e.g.,][]{len07}. Protons, being inevitably accelerated if they co-exist with electrons in the emission regions, probably dominate energetically and dynamically the jets of powerful AGN \\citep[e.g.,][]{cel08}. Applying the synchrotron-proton blazar model \\citep{muc03,rei04} to the quiescent M87 data set yields reasonable agreement model fits that support a highly magnetized compact emission region with approximate equipartition between fields and particles and a total jet power comparable with the above estimates, where protons are accelerated up to $\\sim 10^{9}$ GeV.  Outside of the pc-scale core, the well-known arcsecond-scale jet \\citep[e.g.,][]{bir91,mar02,per05} is also a possible source of IC emission. As both the LAT and TeV telescopes are unable to spatially resolve emission on such small scales, the expected spectral and temporal properties of the predicted emission must be examined. On the observed scales, the dominant seed photon source is the host galaxy starlight, and such an IC/starlight model applied to one of the brightest resolved knots in the jet -- knot A, $\\sim$1 kpc projected distance from the core -- results in a spectrum peaking at TeV energies \\citep{sta05}, thus producing a harder MeV/GeV spectrum than observed by the LAT. Even closer to the core ($\\sim$60 pc, projected), the superluminal knot HST-1 \\citep[$>4c-6c$;][]{bir99,che07} is a more complex case. This knot is more compact than knot A, and its IC emission is expected to be further enhanced by the increased energy densities of the surrounding circumnuclear and galactic photon fields, as well of the comoving synchrotron radiation \\citep{sta06}. The radio/optical/X-ray fluxes of HST-1 have been declining since its giant flare peaked in 2005 \\citep{har09}, with current X-ray fluxes comparable to its pre-flare levels in 2002. Considering the variable and compact nature of the source (with observed months doubling timescales implying $r\\simlt 22 \\delta$ mpc; cf., footnote~\\ref{footnote-hst1}), the predicted IC spectrum has a complex temporal and spectral behavior. In the absence of detailed contemporaneous measurements, its possible role in the production of the LAT observed MeV/GeV emission is unclear.  Continued LAT monitoring of M87 coordinated with multi-wavelength observations can extend the current study of `quiescent' emission to possible flaring, in order to further address the physics of the radiation zone. While the extragalactic $\\gamma$-ray sky is dominated by blazars \\citep{har99,lbas}, this optimistically indicates an emerging population of $\\gamma$-ray radio galaxies.  Other examples, including the few possible associations with EGRET detections like, NGC~6251 \\citep{muk02} and 3C~111 \\citep{sgu05,har08} await confirmation with the LAT, and more radio galaxies are expected to be detected at lower fluxes. This holds great promise for systematic studies of relativistic jets with a range of viewing geometries in the high energy $\\gamma$-ray window opened up by the $Fermi$-LAT."
    },
    "0910/0910.1083_arXiv.txt": {
        "abstract": "{We studied the role of fundamental constants in an updated recombination scenario, focusing on the time variation of the fine structure constant $\\alpha$ and the electron mass $m_e$ in the early Universe. Using CMB data including WMAP 5-yr release, and the 2dFGRS power spectrum, we put bounds on variations of these constants, when both constants are allowed to vary, and in the case that only one of them is variable. In particular, we have found that $-0.019 < \\Delta \\alpha / \\alpha_0 < 0.017$ (95\\% c.l.), in our joint estimation of $\\alpha$ and cosmological parameters.  Finally, we analyze how the constraints depend on the recombination scenario.} ",
        "introduction": "\\label{Intro}  Time variation of fundamental constants is a prediction of theories that attempt to unify the four interactions in nature. Many observational and experimental efforts have been made to put constraints on such variations. Cosmic microwave background radiation (CMB) is one of the most powerful tools to study the early universe and in particular, to put bounds on possible variations of the fundamental constants between early times and the present.  Previous analysis of CMB data (earlier than the WMAP five-year release) including a possible variation of $\\alpha$ have been performed by \\citet{Martins02,Rocha03,ichi06,Stefanescu07,Mosquera07,landau08} and including a possible variation of $m_e$ have been performed by \\citet{ichi06,Scoccola07,landau08,YS03}.   In the last years, the process of recombination has been revised in great detail \\citep{chluba1,chluba2,Hirata09}, and in particular, helium recombination has been calculated very precisely, revealing the importance of considering new physical processes in the calculation of the recombination history \\citep{Dubrovich05,SH08a,SH08b,SH08c,KIV07}.    In a previous paper \\citep{Scoccola08}, we have analized the variation of $\\alpha$ and $m_e$ in an improved recombination scenario. Moreover, we have put bounds on the possible variation of these constants using CMB data and the power spectrum of the 2dFGRS.  In Section \\ref{dependencies}, we review the dependences on $\\alpha$ and $m_e$ in the recombination scenario.  In section \\ref{resultados} we present bounds on the possible variation of $\\alpha$ and $m_e$ using the 5-yr data release of WMAP \\citep{Hinshaw09} together with other CMB experiments and the power spectrum of the 2dFGRS \\citep{2dF05}. In addition to the results of our previous work \\citep{Scoccola08}, here we also present bounds considering only variation of one fundamental constant ($\\alpha$ or $m_e$) alone.  A comparison with similar analyses by other authors is presented in Section \\ref{conclusion}. ",
        "conclusions": "\\label{conclusion}    The obtained results for the cosmological parameters are in agreement within $1 \\sigma$ with the ones obtained by the WMAP collaboration \\citep{wmap5}, without considering variation of fundamental constants. It is also interesting to compare our results with the works of \\citet{Nakashima08} and \\citet{Menegoni09} where only the variation of $\\alpha$ was analized using WMAP 5 year release and the Hubble Space Telescope (HST) prior on the $H_0$. In these works, the HST prior on $H_0$ is used to reduce the large degeneration between $H_0$ and $\\alpha$ and find more stringent constraints on $\\alpha$ variation. However, we have shown in \\citet{Mosquera07} (where the variation of $\\alpha$ was analysed using the WMAP 3 year release), that more stringent constraints on $\\alpha$ can be found using the power spectrum of the 2dFGRS. Indeed, our constraints on $\\alpha$ alone are more stringent than those reported by \\citet{Nakashima08}. On the other hand, our constraints are of the same order than those presented by \\citet{Menegoni09}, using data from higher multipoles reported by recent CMB experiments, such as QUAD \\citep{quad} and BICEP \\citep{bicep}. Finally, it is important to stress, that when using the HST prior on $H_0$, the correct value to be used is the value obtained using only the closest objects, since bounds obtained using other objects could be affected by a possible $\\alpha$ variation."
    },
    "0910/0910.2049_arXiv.txt": {
        "abstract": "{} {The millisecond pulsar PSR J$1740-5340$ in the globular cluster NGC\\,6397 shows radio eclipses over $\\sim40\\%$ of its binary orbit. A first Chandra observation revealed indications for the X-ray flux being orbit dependent as well. In this work, we analysed five datasets of archival Chandra data taken between 2000 and 2007 to investigate the emission across the pulsar's binary orbit.} {Utilizing archival Chandra observations of PSR J1740--5340, we performed a systematic timing and spectral analysis of this binary system.} {Using a $\\chi^{2}$-test, the significance for intra-binary orbital modulation was found to be between 88.5\\% and 99.6\\%, depending on the number of phase bins used to construct the light curve. Applying the unbiased statistical Kolmogorov-Smirnov (KS) test did not indicate any significant intra-binary orbital modulation. However, comparing the counting rates observed at different epochs, a flux variability on times scales of days to years is indicated. The possible origin of the X-ray emission is discussed in a number of different scenarios.} {} ",
        "introduction": "The millisecond pulsar (MSP) PSR J1740--5340 (6397A), located in the core-collapsed globular cluster NGC\\,6397 at a distance of 2.5 kpc, was discovered during a systematic search for ms-pulsars in Galactic globular clusters using the Parkes Radio Telescope \\citep{damico2001}. It is orbiting around a massive late type companion \\citep[$>$ 0.14 $M_{\\odot}$;][]{2003orosz} with an orbital period of 1.35 day \\citep{damico2001b}. The inclination of the binary system to the line of sight is $i\\sim 50^\\circ$ \\citep{2003orosz}. The pulsar's spin period and period derivative are $P=3.65$ ms \\citep{damico2001} and $\\dot{P} = 4.0 \\times 10^{-20} ~\\mathrm{s ~s^{-1}}$ \\citep{2005possenti}. The contribution to $\\dot{P}$  due to possible pulsar acceleration in the cluster's gravitational field is estimated to be smaller than $10^{-20}~\\mathrm{s ~s^{-1}} $  by its large offset of 0.6 arcmin (corresponding to eleven core radii) from the cluster center \\citep{damico2001b}. The spin parameters imply a spin-down luminosity, a characteristic age and a dipole surface magnetic field of \\PSR\\ of $\\dot{E} \\sim 3.3 \\times 10^{34} ~\\mathrm{ergs ~s^{-1}}$, $\\tau_{c} = P/2\\dot{P} \\sim 1.4 \\times 10^{9}$ years, and $B_\\perp \\sim 3.9 \\times 10^{8}$ G, respectively \\citep{2005possenti}.  The radio emission from \\PSR\\ is partially eclipsed in the orbital phase interval $0.05-0.45$ for approximately 40\\% of each orbit \\citep{damico2001,damico2001b}. Moreover, strong fluctuations of the radio signals were observed at nearly all orbital phases \\citep{damico2001b}, which led to the interpretation that the pulsar is orbiting within an extended envelope of matter released from the companion. The optical light curve of the companion showing tidal distortions provides strong evidence that \\PSR\\ is orbiting a companion whose Roche lobe is nearly completely filled \\citep{2001ferraro}. However, the system is in a radio-ejection phase in which accretion is inhibited by the radiation pressure exerted by the pulsar on the infalling matter. The strong interaction between the MSP flux and the plasma wind would explain the irregularities seen in the radio signals from \\PSR\\, \\citep{2001ferraro,2002burderi}.  \\PSR\\ is in a wide orbit with a separation of $\\sim 6.5\\,R_{\\odot}$, so that the wind energy density impinging on the non-degenerate companion should be significantly less than that estimated for other much tighter eclipsing binary systems, such as the field ms-pulsar PSR B1957+20 which has a separation of $\\sim 0.04~R_{\\odot}$ from its companion \\citep{1994arzoumanian}. \\PSR\\ therefore is unlikely to drive a wind of sufficient density off its companion \\citep{damico2001b}. \\citet{2003orosz} have studied the optical light curve of the companion and found no evidence of heating from the pulsar radiation, which supports the aforementioned inference. \\citet{2003sabbi} have investigated the chemical composition of the non-degenerate companion star and found a strong depletion of carbon. This suggests a scenario in which the companion is an evolved star that has lost most of its surface layers. In view of the high rotational velocity of the companion \\citep[$\\sim50$~km/s;][]{2003sabbi}, the stellar wind possibly can be strong enough to cause the mass-loss and result in an extended envelope of matter.  X-ray emission from this binary system was detected with the Chandra X-ray Observatory \\citep{2001grindlay}. Based on an observation in 2000 (cf. Table 1) which only covered $\\sim 40\\%$ of the binary orbit, \\citet{2002grindlay} reported evidence that the count rate of the system appears to increase by a factor of $\\sim2$ at phase 0.4, just before the pulsar comes out of the radio eclipse. However, the short exposure time and the limited orbital phase coverage of this first observation (cf. top panel in Figure~1 and Table~1) did not support a detailed temporal and spectral analysis of the emission.  In this publication we report on a search for X-ray orbital modulation as well as on a spectral analysis of the PSR J1740--5340 binary system making use of archival Chandra data that cover various orbital phases ranges. ",
        "conclusions": "We have searched for the orbital modulation of the X-ray emission from \\PSR\\ associated with NGC 6397. Only one of five archival datasets cover a full binary orbit. Analysing this data with a Kolmogorov-Smirnov test revealed no significant intra-orbital flux modulation. However, correlating the ACIS-S vignetting-corrected net counting rates observed at various orbital phases in 2000, 2002 and 2007 and assuming that the system in general shows no intra-orbital flux modulation reveals a $\\sim 3\\sigma$ flux variability on time scales of days to years (cf.~Col.~7 in Table 1). Figure~\\ref{long} depicts the ACIS-S net count rates observed in the various observations.  \\begin{figure} \\includegraphics[angle=-90,width=9cm]{longTermVariability.ps} \\caption{ACIS-S3 count rates of the \\PSR\\ binary system vs. observation dates.} \\label{long} \\end{figure}  In the following, we will discuss three scenarios which we believe could be the main source of the observed X-rays in the \\PSR\\ binary system.  \\subsection*{1. The Companion star} The companion is observed to be a late-type star. It has an unusual position in the color-magnitude diagram (CMD) as being as luminous as a main-sequence turn-off star but redder \\citep{2001ferraro,2003orosz}. This makes it quite similar to the ``red straggler'' active binaries which are identified to be X-ray sources in other globular clusters \\citep{2001albrow,2003edmonds1,2003edmonds2,2008bassa}. Detailed studies of X-ray sources in various globular clusters by \\cite{2004bassa} and \\cite{2008verbunt} have shown that the separatrix $\\log L_{\\rm 0.5-2.5 keV}=34.0-0.4M_{\\rm V}$ allows one to distinguish between the population of cataclysmic variables and that of magnetically active binaries. Here, $M_{\\rm V}$ is the absolute V-band magnitude, which for the PSR J$1740-5340$ binary system was observed to be $M_{\\rm V} \\sim 5$ \\citep{2003orosz}. Taking this together makes the properties observed from the companion star consistent with that from rapid rotating stars. The X-ray luminosities of late-type stars (i.e. F7 to M5) are found to be well correlated with their equatorial rotational velocities $v_{\\rm rot}$, according to $L_{\\mathrm{X}}\\simeq10^{27}v^{2}_{\\rm rot}$ ergs~s$^{-1}$ \\citep{1981pallavicini}. Adopting the $v_{\\rm rot}\\sim50$~km/s as reported by \\citet{2003sabbi} yields an expected X-ray luminosity of $\\sim2.5 \\times10^{30}$~ergs~s$^{-1}$. Given that the $L_{\\mathrm{X}}$ vs. $v^{2}_{\\rm rot}$ correlation has a scatter of one order of magnitude \\citep{1981pallavicini} it cannot be excluded that the X-ray emission observed from the \\PSR\\ binary is coronal emission from the companion star only. No significant X-ray intra-orbital modulation is observed, which is unlike what we have seen in the optical bands. The optical emission is expected to come from the surface of the companion star while the X-ray emission comes from the corona. As the companion star is interpreted to be tidal-distorted \\citep{2003orosz}, the emitting area responsible for the optical intensity appears to be modulated. However, the X-ray emission is coronal and may be less or not affected by this distortion.  \\subsection*{2. The millisecond pulsar} Considering a possible contribution from the pulsar magnetosphere is justified by the pulsar's spin-down energy of $\\dot{E} \\sim 3.3 \\times10^{34} ~\\mathrm{ergs~s^{-1}}$. By assuming the best-fit X-ray conversion efficiency found for other field and millisecond pulsars, $L_\\mathrm{X} \\sim 10^{-3}-10^{-4} \\dot{E}$ \\citep{2009becker}, the spin-down power of \\PSR\\ implies an X-ray luminosity which is of the order of $\\sim 10^{30} - 10^{31}~\\mathrm{ergs ~s^{-1}}$, in agreement with the observed luminosity (cf.~Table 3). Observing the pulsar in a mode suitable to resolve its X-ray pulses would be worthwhile and could constrain this interpretation.  \\subsection*{3. Intrabinary shock emission} \\citet{2002grindlay} suggested that the X-ray emission they detected from the \\PSR\\ binary system is due to the interaction between the pulsar and shocked gas lifted from the stellar companion. In the shock, the extended matter in the orbit is compressed and gives rise to a power law distribution of synchrotron-emitting particles, $N(\\gamma)\\propto\\gamma^{-p}$, by accelerating a fraction of these particles to very high energies, where $\\gamma$ is the Lorentz factor of the emitting particles. This should lead to an enhancement of X-ray emission from the system as the pulsar enters the eclipsing region and a subsequent decrease due to the increased absorption in this region. \\citet{cheng2006} found that the X-ray luminosity of intrabinary shock emission in the 2--10~keV band can be described by $L^{\\rm th}_{X}\\sim (0.3-1.2)\\times10^{32}D_{11}^{-(p-1)/2}\\dot{E}_{35}^{(p+5)/4}$~ergs~s$^{-1}$, assuming a Lorentz factor for the synchrotron-emitting particles of $\\gamma=(0.5-1.0)\\times10^{6}$, the fractional energy density of electrons to be $\\epsilon_{e}=0.5$ and the fractional energy density of the magnetic field to be $\\epsilon_{B}=0.02-0.1$. $p$ and $D_{11}$ denote the index of the postshock electron energy distribution $N(\\gamma)\\propto\\gamma^{-p}$, and the orbital separation in units of $10^{11}$~cm. The X-ray photon index inferred from the spectrum is $\\Gamma\\sim1.8$ which would suggest that the emission from the shock is in a slow cooling regime \\citep{cheng2006}. The photon index is related to $p$ as $\\Gamma=(p+1)/2$. The observed $\\Gamma$ implies that $p$ is $\\sim2.6$. With $\\dot{E}_{35}=0.33$ and $D_{11}=4.5$, $L^{\\rm th}_{X}$ is found to be in the range of $\\sim (0.1-4.4) \\times10^{30}$~ergs~s$^{-1}$, which is at the observed level. Intrabinary shock emission thus can explain the observed X-ray luminosity though the modulation depth and amplitude, which is also a function of the inclination angle, is unconstrained from this model.  Based on the preceding discussion, it is not possible to firmly distinguish the possible emission scenarios which may, alone or in combination, be responsible for the observed X-ray emission. To disentangle them would require detailed orbital-phase resolved spectroscopy along with a search for X-ray pulses from the millisecond pulsar. Further dedicated X-ray observations covering more binary orbits are certainly needed to further constrain the emission processes at work in this binary system."
    },
    "0910/0910.4258_arXiv.txt": {
        "abstract": "{We present results of the {\\it Suzaku} observation of the dipping, periodically bursting low mass X-ray binary XB\\th 1323-619 in which we concentrate of the spectral evolution in dipping in the energy range 0.8 - 70 keV. It is shown that spectral evolution in dipping is well-described by absorption on the bulge in the outer accretion disk of two continuum components: emission of the neutron  star plus the dominant, extended Comptonized emission of the accretion disk corona (ADC). This model is further supported by detection of a relatively small, energy-independent decrease of flux above 20 keV due to Thomson scattering. It is shown that this is consistent with the electron scattering expected of the bulge plasma. We address the recent proposal that the dip sources may be explained by an ionized absorber model giving a number of physical arguments against this model. In particular, that model is inconsistent with the extended nature of the ADC for which the evidence is now overwhelming. ",
        "introduction": "The dipping sources form a class of about 10 low mass X-ray binaries (LMXB) exhibiting reductions in X-ray intensity at the orbital period; dipping may be shallow or may be 100\\% deep in the band 1 - 10 keV. Dipping in general leads to a hardening of the spectrum due to removal of lower energy photons indicating that it is caused by photoelectric absorption. It has been generally accepted since the 1980s that this takes place in the bulge in the outer accretion disk where the accretion flow from the companion impacts (White \\& Swank 1982; Walter et al. 1982). This plasma has a low ionization state as situated far from the neutron star and so spectral modelling of dipping has employed neutral absorber cross sections as these do not differ substantially from those of low ionization state ions. Sources with more unusual spectral evolution have been investigated, such as 4U\\th 1755-338 in which dipping is approximately energy independent (Church \\& Ba\\l uci\\'nska-Church 1993) and X\\th 1624-490 (Church \\& Ba\\l uci\\'nska-Church 1995) in which a spectral softening takes place. However, the spectral evolution in these sources was shown to be still well-described by photoelectric absorption by cool material providing there are two contiuum emission components in the sources, i.e. blackbody emission from the neutron star, plus Comptonized emission of an accretion disk corona (ADC).  {\\it ASCA} and {\\it BeppoSAX} observations of the dipping sources XB\\th 1916-053 led to proposal of the ``progressive covering'' explanation of dipping (Church et al. 1997, 1998a,b). The spectral evolution in dipping in this source and dip sources in general is relatively complex with parts of the spectrum being absorbed while other parts at lower energies were not absorbed. This may be explained in terms of the major Comptonized emission component being extended such that it is gradually overlapped by the extended absorber so that in any stage of dipping the part of the extended ADC emission covered by the bulge is absorbed producing the absorbed part of the spectrum, while the uncovered part is not covered producing the unabsorbed part of the spectrum. As dipping proceeds, this part becomes a smaller fraction of the total as observed and as confirmed by spectral fitting. Spectral modelling on this basis gave very good descriptions of dipping, for example, in the {\\it ASCA} observations of XB\\th 1916-053 (Church et al. 1997) and of XBT\\th 0748-676 (Church et al. 1998a) and of the {\\it BeppoSAX} observations of XB\\th 1916-053 (Church et al. 1998b). It has since been shown to apply to the dipping sources generally.  The extended nature of the Comptonizing ADC has since then been demonstrated in the dipping sources using the technique of dip ingress timing. The time taken for the transition from non-dip to 100\\% deep dipping $\\Delta t$ is proportional to the size of the extended emitter: $r_{\\rm ADC}$ = $\\pi \\, r_{\\rm AD}\\, \\Delta t/ P$, where $r_{\\rm ADC}$ is the radial extent of the corona, $r_{\\rm AD}$ is the radial extent of the accretion disk and $P$ is the orbital period. Application of this technique to the dipping LMXB (Church \\& Ba\\l uci\\'nska-Church 2004) gave ADC radial sizes as shown in Fig. 1 between 20\\th 000 and 700\\th 000 km, depending on source luminosity so that for a bright source ($L$ $\\sim$ $1\\times 10^{38}$ erg s$^{-1}$), the size is $\\sim$500\\th 000 km. \\begin{figure}[!ht]                                                   % \\includegraphics[width=44mm,height=64mm,angle=270]{balucinska-church_2009_01_fig01.ps} % \\caption{Measured values of the radial extent of the ADC in LMXB; from Church \\& Ba\\l uci\\'nska-Church (2004).} \\end{figure}  Recently this has been confirmed by an independent technique based on {\\it Chandra} data by Schulz et al. (2009). Precise grating measurements of the spectrum of Cygnus\\th X-2 reveal a wealth of emission lines of highly excited states such as the H-like ions of Ne, Mg, Si, S and Fe. The width of these lines indicated Doppler shifts due to orbital motion in the accretion disk corona at radial positions between 18\\th 000 and 110\\th 000 km in good agreement with the overall ADC size from dip ingress timing. The evidence for the extended corona is thus now overwhelming.  \\subsection{The ionized absorber model}  In recent years, high resolution instruments on {\\it XMM} and {\\it Chandra} have revealed the presence of absorption lines of highly ionized species, for example in XBT\\th 0748-676 (Jimenez-Garate et al. (2003), and in X\\th 1624-490 (Parmar et al. 2002). However, this proof of the existence of highly ionized states was taken much further by Diaz-Trigo et al. (2006) and Boirin et al. (2005) who proposed that \\hbox{X-ray} \\begin{figure*}[!ht]                                                   % \\includegraphics[width=44mm,height=134mm,angle=270]{balucinska-church_2009_01_fig02.ps} % \\caption{XIS lightcurve of the {\\it Suzaku} observation of XB\\th 1323-619 in the energy band 0.2 - 12 keV with 64 s binning.} \\end{figure*} dipping in all LMXB dipping sources could be explained by an ionized absorber model, i.e. showing that the spectra can be fitted with this model. Fitting the spectral evolution in dipping with this model implies an absorbing medium of high and variable ionization state (log$\\,\\xi$ $\\sim$3), which therefore cannot be in the bulge in the outer disk but must be located much closer to the neutron star. In dipping log$\\,\\xi$ decreases from 3.5 to 3.0 while the column density of ionized absorber increases $\\sim$10 times, with very little change in the column density of neutral absorber.  However, the nature of the absorbing structure is not specified and the reason why the bulge in the outer disk does not absorb is not addressed. The implications of this model, if correct, would be far-reaching for our understanding of LMXB. The proponents of the model argue that their fitting does not require progressive covering, and therefore the model does not require there to be an extended corona. Conversely, the evidence for the extended corona is now so strong that dipping cannot correctly be modelled without progressive covering.  In this work we present results of a recent observation of the dipping source XB\\th 1323-619 with {\\it Suzaku}, viewing the results in the context of cool {\\it versus} ionized absorber. ",
        "conclusions": "There are strong objections to the ionized absorber model as detailed above. The standard explanation as absorption in the bulge does not have these objections and is supported by the present work. The present work supports the extended nature of the Comptonizing ADC for which there is now overwhelming evidence."
    },
    "0910/0910.1361_arXiv.txt": {
        "abstract": "A convincing detection of primordial non-Gaussianity in the local form of the bispectrum, whose amplitude is given by the $\\fnl$ parameter, offers a powerful test of inflation. In this paper we calculate the modification of  two-point cross-correlation statistics of weak lensing - galaxy-galaxy lensing and galaxy-Cosmic Microwave Background (CMB) cross-correlation - due to $\\fnl$. We derive and calculate the covariance matrix of galaxy-galaxy lensing including cosmic variance terms.  We focus on large scales ($l<100$) for which the shape noise of the shear measurement becomes irrelevant and cosmic variance dominates the error budget. For a modest degree of non-Gaussianity, $\\fnl=\\pm 50$, modifications of the galaxy-galaxy lensing signal at the 10\\% level are seen on scales $R\\sim 300$~Mpc, and grow rapidly toward larger scales as $\\propto R^2$.  We also see a clear signature of the baryonic acoustic oscillation feature in the matter power spectrum at $\\sim 150$~Mpc, which can be measured by next-generation lensing experiments. In addition we can probe the local-form primordial non-Gaussianity in the  galaxy-CMB lensing signal by correlating the lensing potential reconstructed from CMB with high-$z$ galaxies. For example, for $\\fnl=\\pm 50$, we find that the galaxy-CMB lensing cross power spectrum is modified by $\\sim 10$\\% at $l\\sim 40$, and by a factor of two at $l\\sim 10$, for a population of galaxies at $z=2$ with a bias of 2. The effect is greater for more highly biased populations at larger $z$; thus, high-$z$ galaxy surveys cross-correlated with CMB offer a yet another probe of primordial non-Gaussianity. ",
        "introduction": "Why study non-Gaussianity? For many years it was recognized that the simple inflationary models based upon a single slowly-rolling scalar field would predict nearly Gaussian primordial fluctuations. In particular, when we parametrize the magnitude of non-Gaussianity in the primordial curvature perturbations $\\zeta$, which gives the observed temperature anisotropy in the Cosmic Microwave Background (CMB) in the Sachs--Wolfe limit as $\\Delta T/T=-\\zeta/5$, using the so-called non-linear parameter $\\fnl$ \\cite{komatsu/spergel:2001} as $\\zeta({\\mathbf x})=\\zeta_L({\\mathbf x})+(3\\fnl/5)\\zeta_L^2({\\mathbf x})$, then the bispectrum of $\\zeta$ is given by\\footnote{Definition of the bispectrum in terms of Fourier coefficients of $\\zeta$ is $\\langle\\zeta_{{\\mathbf k}_1}\\zeta_{{\\mathbf k}_2}\\zeta_{{\\mathbf k}_3}\\rangle=(2\\pi)^3\\delta({\\mathbf k}_1+{\\mathbf k}_2+{\\mathbf k}_3)B_\\zeta(k_1,k_2,k_3)$. Throughout this paper we shall order $k_i$ such that $k_3\\le k_2\\le k_1$.} $B_\\zeta(k_1,k_2,k_3)=(6\\fnl/5)\\left[P_\\zeta(k_1)P_\\zeta(k_2)+(\\mbox{2 cyclic terms})\\right]$, where $P_\\zeta(k)\\propto k^{n_s-4}$ is the power spectrum of $\\zeta$ and $n_s$ is the tilt of the power spectrum, constrained as $n_s=0.960\\pm 0.013$ by the WMAP 5-year data \\cite{komatsu/etal:2009}. This form of the bispectrum has the maximum signal in the so-called squeezed triangle for which $k_3\\ll k_2\\approx k_1$ \\cite{babich/creminelli/zaldarriaga:2004}. In this limit we obtain \\begin{equation}\\label{eq:1} B_\\zeta(k_1,k_1,k_3\\to 0)=\\frac{12}5\\fnl P_\\zeta(k_1)P_\\zeta(k_3). \\end{equation} The earlier calculations showed that $\\fnl$ from single-field slow-roll inflation would be of order the slow-roll parameter, $\\epsilon\\sim 10^{-2}$ \\cite{salopek/bond:1990,falk/rangarajan/srendnicki:1993,gangui/etal:1994}. However, it is not until recent that it is finally realized that the coefficient of $P_\\zeta(k_1)P_\\zeta(k_3)$ from the simplest single-field slow-roll inflation with the canonical kinetic term in the squeezed limit is given precisely by \\cite{maldacena:2003,acquaviva/etal:2003} \\begin{equation} B_\\zeta(k_1,k_1,k_3\\to 0)=(1-n_s) P_\\zeta(k_1)P_\\zeta(k_3). \\label{eq:singleprediction} \\end{equation} Comparing this result with the form predicted by the $\\fnl$ model, one obtains $\\fnl=(5/12)(1-n_s)$.  Perhaps, the most important theoretical discovery regarding primordial non-Gaussianity from inflation over the last few years is that, not only models with the canonical kinetic term, but {\\it all} single-inflation models predict the bispectrum in the squeezed limit  given by Eq.~(\\ref{eq:singleprediction}), regardless of the form of potential, kinetic term, slow-roll, or initial vacuum state \\cite{creminelli/zaldarriaga:2004,seery/lidsey:2005,chen/etal:2007,cheung/etal:2008}. Therefore, the prediction from all single-field inflation models is $\\fnl=(5/12)(1-n_s)= 0.017$ for $n_s=0.96$. A convincing detection of $\\fnl$ well above this level is a breakthrough in our understanding of the physics of very early universe \\cite{bartolo/etal:2004,komatsu/etal:prepb}. The current limit from the WMAP 5-year data is $\\fnl=38\\pm 21$ (68\\%~CL) \\cite{smith/senatore/zaldarriaga:prep}.   There are many ways of measuring $\\fnl$. The most popular method has been the bispectrum of CMB \\cite{verde/etal:2000,wang/kamionkowski:2000,komatsu/spergel:2001,komatsu/etal:2002,komatsu/etal:2003} (also see \\cite{komatsu:prep} for a pedagogical review). The other methods include the trispectrum of CMB \\cite{okamoto/hu:2002,kogo/komatsu:2006}, the bispectrum of galaxies \\cite{scoccimarro/etal:2004,sefusatti/komatsu:2007,jeong/komatsu:2009b,sefusatti:prep}, and the abundance of galaxies and clusters of galaxies \\cite{lucchin/matarrese:1988,matarrese/verde/jimenez:2000,sefusatti/etal:2007,loverde/etal:2008}.  Recently, analytical \\citep{dalal/etal:2008,matarrese/verde:2008,slosar/etal:2008,afshordi/tolley:2008,taruya/etal:2008} and numerical \\citep{dalal/etal:2008,desjacques/seljak/iliev:prep,pillepich/porciani/hahn:prep,grossi/etal:prep} studies of the effects of primordial non-Gaussianity on the power spectrum of dark matter halos, $P_h(k)$, have revealed an unexpected signature of primordial non-Gaussianity in the form of a {\\it scale-dependent galaxy bias}, i.e., $P_h(k)=b_1^2P_m(k)\\to [b_1+\\Delta b(k)]^2P_m(k)$, where $P_m(k)$ is the power spectrum of matter density fluctuations, and \\begin{equation} \\Delta b(k) = \\frac{3(b_1-1)\\fnl\\Omega_mH_0^2\\delta_c}{D(z)k^2T(k)}. \\label{eq:bk} \\end{equation} Here, $D(z)$ and $T(k)$ are the growth rate and the transfer function for linear matter density fluctuations, respectively, and $\\delta_c=1.68$ is the threshold linear density contrast for a spherical collapse of an overdensity region. The $k^2$ factor in the denominator of $\\Delta b(k)$ shows that this effect is important only on very large scales. Highly biased tracers are more sensitive to $f_{NL}$. ",
        "conclusions": "In this paper we have studied the galaxy-galaxy lensing and galaxy-CMB lensing cross-correlation functions. We have focused on large scales, typically larger than 100~Mpc at the lens redshift. While current measurements have high signal-to-noise ratios on much smaller scales, we believe that future surveys will enable detection of interesting physical effects in the large-scale, linear regime.  We derive the full covariance matrix for galaxy-galaxy lensing, including the cosmic variance due to the clustering of lenses and to cosmic shear (Eq.~\\ref{eq:cov}). We use the linear bias model to provide the halo-mass and halo-halo correlations needed for this calculation. We present results for the covariance of the mean tangential shear measurement as a function of angular separations, as well as for the harmonic space halo-convergence cross-power spectrum. Our calculations show that the errors in $\\Delta\\Sigma(R)$ are dominated by the cosmic variance term for $R\\gtrsim 50~h^{-1}~{\\rm Mpc}$ (see Fig.~\\ref{fig:bao_ebar}). Similarly, the errors in the halo-convergence cross power spectra, $C_l^{h\\kappa}$, are dominated by the cosmic variance term at $l\\lesssim 100$ (see Fig.~\\ref{fig:cl_cmb_lowz_err}).  For Gaussian initial conditions, we show that the baryonic effects in the matter power spectrum (often called Baryon Acoustic Oscillations) produce a ``shoulder'' in the galaxy-galaxy lensing correlation (i.e., the mean tangential shears), $\\Delta\\Sigma(R)$, at $R\\sim 110~h^{-1}~{\\rm Mpc}$ (see Fig.~\\ref{fig:bao}). This effect should be easy to measure from the next-generation lensing surveys by combining $\\Delta\\Sigma(R)$ from multiple lens redshift slices.  We consider the prospects of detecting primordial non-Gaussianity of the local-form, characterized by the $\\fnl$ parameter. We have found that the scale-dependent bias from the local-form non-Gaussianity with $\\fnl=\\pm 50$ modifies $\\Delta\\Sigma(R)$ at the level of 10--20\\% at $R\\sim 300~h^{-1}~{\\rm Mpc}$ (depending on $b_1$ and $z_L$; see Fig.~\\ref{fig:fnl_gglens_diff}) (see Fig.~\\ref{fig:fnl_gglens_signal}). The modification grows rapidly toward larger scales, in proportion to $R^2$. High-$z$ galaxies at, e.g., $z\\gtrsim 2$, cross-correlated with CMB can be used to find the effects of $\\fnl$ in the galaxy-convergence power spectrum, $C_l^{h\\kappa}$. While the effects are probably too small to see from a single lens redshift (see Fig.~\\ref{fig:cl_cmb_highz_err}), many slices can be combined to beat down the cosmic variance errors. Exactly how many slices are necessary, or what is the optimal strategy to measure $\\fnl$ from the galaxy-CMB lensing signal requires a more detailed study that incorporates the survey strategy for specific galaxy and lens surveys.  We emphasize that, while the two-point statistics of shear fields are not sensitive to primordial non-Gaussianity, the two-point statistics correlating shear fields with density peaks (i.e., galaxies and clusters of galaxies) are sensitive due to the strong scale-dependence of halo bias on large scales.  Finally, we note that one can also measure the effects of $\\fnl$ on the halo power spectrum, $C_l^h$. For example, $C_l^h$ that would be measured from LSST can be used to probe $\\fnl\\sim 1$ \\citep{carbone/verde/matarrese:2008}; thus, we would expect $C_l^h$ to be more powerful than the lens cross-correlation statistics we studied here. However, a combination of the two measurements would provide useful cross-checks, as galaxy clustering and galaxy-lensing correlations are affected by very different systematics."
    },
    "0910/0910.0367_arXiv.txt": {
        "abstract": "A detailed analysis of a long \\textit{XMM-Newton} observation of the Narrow Line Seyfert 1 galaxy 1H0707-495 is presented, including spectral fitting, spectral variability and timing studies. The two main features in the spectrum are the drop at $\\sim7$ \\keV and a complex excess below 1 \\keV. These are well described by two broad, K and L, iron lines. Alternative models based on absorption, although they may fit the high energy drop, cannot account for the 1 \\keV complexity and the spectrum as a whole. Spectral variability shows that the spectrum is composed of at least two components, which are interpreted as a power-law dominating between 1--4 \\keV, and a reflection component outside this range. The high count rate at the iron L energies has enabled us to measure a significant soft lag of $\\sim 30 $ s between 0.3--1 and 1--4 \\keV, meaning that the direct hard emission leads the reflected emissions. We interpret the lag as a reverberation signal originating within a few gravitational radii of the black hole. ",
        "introduction": "In a recent paper, \\cite{2009Natur.459..540F} reported on the discovery of the first broad iron L line in the Narrow Line Seyfert 1 galaxy 1H0707-495. The interpretation is that the line is broadened by relativistic effects very close to the black hole, where the line is thought to originate.\\\\ This, along with the already established broad iron K line seen in many AGN, is a further confirmation that the observed X-ray spectra from suppermassive black holes in radiatively efficient AGN are emitted from within few gravitational radii of the black hole (\\citealt{1995Natur.375..659T,2007MNRAS.382..194N}, see \\citealt{2007ARA&A..45..441M} for a review).  The standard picture is that matter is accreted in the form of an optically thick, geometrically thin disc (\\citealt{1973A&A....24..337S}) that radiates mainly in the UV. The usually observed power-law spectrum in X-rays, which is likely to be due to a Comptonising corona and/or the base of jet, is expected to illuminate the disc, producing a characteristic reflection spectrum that is observed alongside the primary power-law (\\citealt{1991MNRAS.249..352G}). The main feature in the reflection spectrum is a fluorescent Fe K$\\alpha$ line at $\\sim 6.4$ \\keV, which is expected to be broadened and smeared out if it is emitted close to the black hole. Other features are also expected to be seen in certain circumstances such as high iron abundance (\\citealt{2009Natur.459..540F}, this work), while a soft excess is naturally produced for a wide range of ionisation states.  Although alternative interpretations of AGN spectra exist (e.g. \\citealt{2009A&ARv..17...47T} for a review), they generally fail to explain some variability properties seen in the data (see the discussion section). Spectral variability and timing analysis provide a completely independent way of studying the physics of emission. That, combined with the spectral fitting, can be used to distinguish between the different models. In particular for the reflection interpretation, the model predicts a time delay, that in principle, is observable between the primary emission and the reflected emission (e.g. \\citealt{2001AdSpR..28..267P}). This has been seen for the first time in 1H0707-495 (\\citealt{2009Natur.459..540F}) and is discussed further in this paper.  1H0707-495 ($z=0.0411$, hereafter 1H0707), is a highly variable Narrow Line Seyfert 1 (NLS1) galaxy (\\citealt{1999ApJS..125..297L}). It has a steep spectrum and shows a strong soft excess, and although these are common among other NLS1, they are extreme in 1H0707 (\\citealt{1999MNRAS.309..113V}). The soft excess is traditionally interpreted as disc blackbody emission. However, when large samples of AGN are considered, the measured temperature is higher than the temperature of a standard disc (\\citealt{1973A&A....24..337S}), and independent of blackhole mass (\\citealt{2004MNRAS.349L...7G}). It appears to be well explained by the blend of relativistically smeared reflection continuum (\\citealt{2006MNRAS.365.1067C}; see \\citealt{2008MNRAS.386L...1S} for problems with alternative absorption models).  1H0707 gained interest after the discovery of a sharp drop in the spectrum at $\\sim 7$ \\keV, during the first XMM-Newton observation [O1 hereafter] (\\citealt{2002MNRAS.329L...1B}). A second XMM observation (O2) showed a change in the energy of the edge from 7.1 to 7.5 \\keV (\\citealt{2004MNRAS.353.1064G}). The drop was attributed to either the source being partially covered (\\citealt{2002MNRAS.329L...1B,2004PASJ...56L...9T}), or it being the blue wing of a broad iron K$\\alpha$ line (\\citealt{2004MNRAS.353.1071F}). In 2007, four short (40 ks each) observations were made with XMM (O3 hereafter). A short report on a 4th much longer observation (O4 hereafter) done in 2008 was presented in \\cite{2009Natur.459..540F}. In this paper, we present detailed analysis of the results reported there. We also include analysis of the O3 data in some parts of this study which is organised as follows: observation and data reduction are presented in Sec. \\ref{observation}, spectral fitting and variability are presented in detail in Sec. \\ref{spec_fitting} and \\ref{spec_var} respectively. In Sec. \\ref{timing}, the lag calculation are shown, with all the results and implications discussed in Sec. \\ref{discussion}. \\begin{figure} \\centering \\includegraphics[height=210pt,angle=270,clip ]{figures/unfolded_spec.ps} \\caption{Unfolded spectrum of 1H0707 in the 2008 XMM observations (O4) compared to the first observation (O1). The spectra have been unfolded (using \\texttt{eeufspec}) against a power-law of index 0. This removes the effect of the detector effective area and shows the main features in the spectrum. O4 represents the spectrum in the `bright' state, and the O1 represents a typical spectrum for the `faint' state.} \\label{fig:spec_all} \\end{figure} ",
        "conclusions": "\\label{discussion} A detailed analysis of a long XMM observation of 1H0707 has been presented. It is clear that the source during the present O4 observation is in a state similar to the earlier O2 reported in \\cite{2004MNRAS.353.1064G}. The high energy spectral drop has an energy of $7.31^{+0.12}_{-0.06}$ \\keV. The two models, partial covering and reflection, both give an equally good fit to the data above 3 \\keV. However, as noted in \\cite{2009MNRAS.tmpL.244R} for MCG-6-3-15, an absorber that is responsible for the spectral drop is expected to produce fluorescent emission at $\\sim 6.4$ \\keV. This is clearly not seen in the data, and might require a special geometry, where the absorber covers only a small part of the sky as seen by the source. The value of $\\sim 20$ per cent covering fraction for the absorber at the source, is an upper limit, and although it is not as small as that reported in \\cite{2009MNRAS.tmpL.244R}, it points to the fact that a `special geometry' explanation cannot hold for all AGN. Since, to reproduce the drop, most of the absorber must lie exactly in our line of sight with a covering fraction of 70 per cent, and a global covering fraction (i.e. as seen by the source) of about 20 per cent.  The drop is more consistent with being the blue wing of a broad K$\\alpha$ emission line, originating very close to the black hole. The inner radius of the emitting region is $\\sim1.3 r_{\\rm g}$, and if we assume no radiation is emitted from within the radius of marginal stability, the black hole has to be an almost maximally-spinning Kerr black hole. A fit using \\texttt{kerrdisc} model from \\cite{2006ApJ...652.1028B} gives a spin parameter $a>0.976$ at 90 per cent confidence (where a is the dimensionless parameter $a=cJ/GM^{2}$, with $J$ being the angular momentum of a black hole of mass $M$, see \\citealt{2009ApJ...697..900M} for spin measurements using this method for stellar mass black holes). Fig. \\ref{fig:spin} shows the $\\chi^2$ changes for the spin parameter. The parameter is well constrained to be $>0.97$, and this is the case also if it is measured from the low flux observations (O1 and the first of O3). Although this might not be an exact value (\\citealt{2006ApJ...652.1028B}, submitted), what is clear from the shape of the line is that most of the radiation has to be emitted from within a few gravitational radii, and the black hole in 1H0707 appears to be rapidly spinning similar to what was found for MCG-6-30-15 (\\citealt{2006ApJ...652.1028B}).  This interpretation is further strengthened with the identification of the feature at $\\sim 1$ \\keV as the a broad iron L-shell line (\\citealt{2009Natur.459..540F}). As we have shown in Sec. \\ref{soft_fit} (see Fig. \\ref{fig:base_plus_diff}), the spectral complexities around 1 \\keV cannot be fitted with any absorption-only model, and there is a requirement for an emission component. The flux ratio between the K and L features is 1:20, measured as the ratio of equivalent widths of two \\texttt{laor} lines, is consistent with the value expected from atomic physics (\\citealt{2005MNRAS.358..211R}).  It is apparent from the models fitted to the high energy band that the spectrum requires high iron abundance. This is expected to imprint other signatures in the spectrum. In particular, if the apparent drop at $\\sim 1$ \\keV is interpreted as an iron L edge, then a deep absorption trough is expected to be produced by iron M-shell UTA. A reflection continuum on the other hand, accounts naturally for most of the features.  The high iron abundance ($\\sim 9\\times$ solar), if real, is not inconsistent with what is expected from line ratio modeling at other wavelengths (\\citealt{2002ApJ...567L..19S,2002ApJ...564..592H}), and would imply that the host galaxy had gone through a phase of strong star formation. This seem to be consistent with studies showing that NLS1 have stronger star formation than normal Seyferts (\\citealt{2009arXiv0908.0280S}). Also the FeII emission in also strong in these sources, and it can also be attributed to high abundance (\\citealt{2000NewAR..44..531C}). The high abundance may also be due to a dense nuclear supernova activity. On the modeling side, non-solar values for elements other than iron can slightly affect the shape of the broad iron K$\\alpha$ and the abundances inferred from the fits (\\citealt{1995MNRAS.276.1311R}). \\begin{figure} \\centering \\includegraphics[width=180pt ]{figures/spin.eps} \\caption{$\\Delta\\chi^2$ versus spin parameter for a fit using \\texttt{kerrdisc}. The black hole is rapidly spinning. The fit was made to 3.0-10.0 \\keV band for the combined O4 observations. \\textit{Inset:} a zoom-in on the high spin values. The dashed line in both plots marks the $3\\sigma$ limit.} \\label{fig:spin} \\end{figure}  Because the source is highly variable, fitting a time-averaged spectrum might not give the whole picture, Both flux and time-resolved spectroscopy indicate that the spectrum is composed of two components, a highly variable component between 1--4 \\keV and a less variable component outside the range. These two components are consistent with a PLC and RDC respectively, similar to the results from other AGN (\\citealt{2003MNRAS.340L..28F,2003MNRAS.342L..31T,2004MNRAS.348.1415V}). In particular, the shape of the constant (or slowly varying) component dominating the low flux state matches very well the shape of a reflection spectrum (see Fig. \\ref{fig:spec_all}). Variations in the power-law component drives most of of the variability (Fig. \\ref{fig:pow_ref_all}), and this simple model reproduces the variability pattern very well (Fig. \\ref{fig:flux_resolved_image}).  No absorption-only model provides a good fit to the spectrum, and these models can also be ruled out based on spectral variability. If the spectral drop at $\\sim 7$ \\keV is due to an absorbing cloud in the line of sight, then the depth and shape of the drop imply the absorber only partially covers the source, which in turn means it has to be separate from any absorber causing the drop at $\\sim 1$ \\keV. This however, is inconsistent with the variability pattern showing that the spectral features at $\\sim 7$ and $\\sim 1$ \\keV are varying in a similar way (Fig. \\ref{fig:flux_resolved_image}).  The variability pattern, where a good correlation between reflection and the direct continuum is seen at low fluxes and levels off at higher fluxes, is consistent with the predictions of the light bending model (\\citealt{2003MNRAS.340L..28F,2004MNRAS.349.1435M}). Most of the luminosity emerges from within few gravitational radii, and strong gravitational light bending \\textit{must} be taken into account. However, the model in its simplest form assumes that the intrinsic luminosity of the source is constant, which most likely is not the case. Also, assuming a single ionisation parameter throughout the disc might not be realistic too. These two factors may be responsible for producing the positive correlation in Fig. \\ref{fig:pow_xi_all}, where an increase in the intrinsic luminosity of the primary component causes the disc ionisation to increase. The most likely scenario is that both light bending and intrinsic source variations are at work. \\subsection{Time Lag and Reverberation} The lag spectrum is frequency dependent, and cannot be fitted with a single power-law. If however, the fit is only applied to the central frequency bins, a good fit is found with $\\tau(f)\\propto f^{-0.77}$. The slope cannot be constrained very well, and a power-law with an index of 1 is also consistent with the data. This is similar to (or at least consistent with) what was found in NGC 7469 (\\citealt{2001ApJ...554L.133P}), MCG-6-30-15 (\\citealt{2003MNRAS.339.1237V}), NGC 4051 (\\citealt{2004MNRAS.348..783M}) and NGC 3783 (\\citealt{2005ApJ...635..180M}). Several models have been suggested to explain the lag, the most successful of which is the propagation of accretion fluctuations (\\citealt{1997MNRAS.292..679L,2001MNRAS.327..799K}, see also \\citealt{2006MNRAS.367..801A}). In this model, accretion rate variations are produced in different radii in the accretion disc, each with time-scale associated with the radius of origin. These fluctuations propagate inward and modulate the central X-ray emitting region. This seems to explain most of the timing properties including the PSD shape, its changes with energy and RMS-flux relation (\\citealt{2004MNRAS.347L..61U}). If additionally, energy-dependent emissivity profiles are assumed, with soft energies having flatter emissivity profiles, the dependency of the lag on frequency and energy can also explain the observations in galactic black holes and AGN.  The lag is produced when large time-scale fluctuations propagating inward passes through the soft-emitting region first (larger radii), before reaching the hard-emitting region (\\citealt{2001MNRAS.327..799K}). This produces the positive lag below $\\sim 5\\times10^{-4}$ Hz. The magnitude of the lag itself is dependent on the propagation speed and the distance between the soft- and hard-emitting regions.  The slope of the lag spectrum appears to change at small temporal frequencies (large time-scales). If this turnover is real (i.e. the last point in Fig. \\ref{fig:time_lag}), it might be the effect of the size of the emitting region. The emissivity profiles inferred from the spectral fitting to the iron line indicate that most of the radiation originates within few gravitational radii. This might pose an upper limit on the hard lags that can be observed. Interestingly, the maximum of the lag, at $1-2\\times10^{-4}$ Hz, is about the same as the break frequency in the PSD, which was found to be $1.6^{+0.4}_{-0.6}\\times10^{-4}$ Hz, (when fitted with a broken power-law). The break in the PSD in this propagation model may be related to the time-scale of the fluctuations at the outer radius of the X-ray emitting region (\\citealt{2001MNRAS.327..799K,2004MNRAS.348..783M}). The light crossing distance corresponding to $\\sim150$ s for a black hole mass of $5\\times10^{6}\\rm{M}_\\odot$ is $\\sim 8\\ r_{\\rm g}$. The viscous time-scale at this radius in a standard disc (\\citealt{1973A&A....24..337S}) is $t_{\\rm{visc}}\\sim 10^{6}$ s (assuming a viscosity parameter $\\alpha=0.1$ and a scale height $H/R =0.1$), which is much higher than what is observed. If however, thermal time-scales are considered where $t_{\\rm{th}}=(H/R)^{2}t_{\\rm{visc}}$ (which might be the case if variability is caused by ionisation changes), then the lag peak and the PSD break would correspond to the thermal time-scale of the outer edge of an emitting region that is less than $8 r_{\\rm g}$ in size.  At time-scales below about half an hour, the lag becomes negative, where hard energies now \\textit{lead} softer energies. This is the first significant soft lag detected from black holes (\\citealt{2009Natur.459..540F}). Other authors have reported soft lags before, but they were not significant (e.g. \\citealt{2007MNRAS.382..985M} for Ark 564 and \\citealt{2004MNRAS.347..269G} for IRAS 13224-3809). If the positive lag is interpreted as being due to Comptonisation, where the variability is seen first in the soft band (disc component in the form of reflection) before the hard band (Comptonised coronal emission), then negative soft lag above $\\sim 6\\times10^{-4}$ Hz is simply in the wrong direction. More sophisticated models may be tuned to produce soft lags. In particular, \\cite{2000A&A...359..843M} studied the temporal evolution of flares in a disc-corona system. They found that, if a flare originates in the corona (e.g. by magnetic reconnection) with the perturbation time-scale of the order of few corona light crossing time, then a soft lag is expected, and is caused by a feedback mechanism that cools the corona down. Fluctuations are seen in harder energies, before reprocessed soft photons from the disc cool the corona and softens the spectrum. In this picture, the whole spectrum originates in the corona. This however, is inconsistent with the spectral variability patterns seen in Sec. \\ref{spec_var}. There is an abrupt change in variability properties between the bands above and below $\\sim 1$ \\keV (see also RMS spectrum in Fig. 8 of the supplementary material in \\citealt{2009Natur.459..540F}), which argues for two \\textit{separate} components above and below 1 keV.  In a reflection scenario, the soft lag is naturally expected. The variability is seen first in the power-law dominated component, emitted from the corona. Some of the radiation is reflected off the optically thick disc. This produces a lag comparable to the light travel time between the the corona and the reflector. If we assume most of the radiation is emitted at $\\sim 2 r_{\\rm g}$, implied from the inner radius and emissivity index in the spectral fitting, then interpreting the 30 s lag as the light crossing time yields a black hole mass of $\\rm{M}\\sim2\\times10^6 \\rm{M}_\\odot$, which is consistent with the uncertain mass of this black hole quoted in the literature (e.g. \\citealt{2005ApJ...618L..83Z}).  Ark 564 shows a similar lag spectrum to the one presented here (\\citealt{2007MNRAS.382..985M}), with a difference in the magnitude of the lag. The ratio of hard to soft lags in Ark 564 is $\\sim700/10=70$, compared to $\\sim5$ in 1H0707. This appears to be related to the compactness of the emission region and the strength of the reflection. 1H0707 has a strong reflection component associated with a compact emission region with strong light-bending effects. However Ark 564 has weaker reflection which suggests a more extended emitting region.  Since thermal and viscous time-scales scale as $(r/r_g)^{3/2}$, the hard lags due to changes in the accretion flow in this region are larger compared to the light-travel time.  Unlike other well-studied AGN (e.g. MCG-6-30-15 and Mrk 766), the spectrum of 1H0707 does not show any signs of warm absorption, giving an uninterrupted view to the inner regions of the accretion system. 1H0707 has one of the strongest soft excesses among other NLS1 (e.g. \\citealt{2006MNRAS.368..479G}). That, in our interpretation, is simply because of the high iron abundance making the iron L complex strong, which might also be linked to the conclusion of \\cite{2006MNRAS.368..479G}, that most NLS1 with complex spectral features show high values of Fe~\\textsc{ii}/H$\\beta$. Other sources that are similar to 1H0707 include IRAS 13224-3809, which shows clearly the iron L broad feature, however the current XMM observation is not long enough to check if there is any soft lags (Ponti et al. 2009, submitted).  In summary, although different models may be made to explain some of the properties we see independently, most of them fail to explain all of what we see. In this work, we have explored the data using different methods, and shown that a model in which the contribution of X-ray reflection off optically thick material is significant, provides a natural explanation to the spectrum, spectral variability, as well as the soft lag seen for the first time in 1H0707."
    },
    "0910/0910.2275.txt": {
        "abstract": "A trend is emerging regarding the progenitor stars that give rise to the most common core-collapse supernovae (SNe), those of Type II-Plateau (II-P): they generally appear to be red supergiants with a limited range of initial masses, $\\sim$8--16 M$_{\\sun}$. Here we consider another example, SN 2008cn, in the nearly face-on spiral galaxy NGC~4603.  Even with limited photometric data, it appears that SN 2008cn is not a normal SN II-P, but is of the high-luminosity subclass. Through comparison of pre- and post-explosion images obtained with the Wide Field and Planetary Camera 2 (WFPC2) onboard the {\\sl Hubble Space Telescope (HST)}, we have isolated a supergiant star prior to explosion at nearly the same position as the SN.  We provide evidence that this supergiant may well be the progenitor of the SN, although this identification is not entirely unambiguous. This is exacerbated by the distance to the host galaxy, 33.3 Mpc, making SN 2008cn the most distant SN II-P yet for which an attempt has been made to identify a progenitor star in pre-SN images.  The progenitor candidate has a more yellow color ($[V-I]_0=0.98$ mag and $T_{\\rm eff} = 5200 \\pm 300$ K) than generally would be expected and, if a single star, would require that it exploded during a ``blue loop'' evolutionary phase, which is theoretically not expected to occur. Nonetheless, we estimate an initial mass of $M_{\\rm ini} = 15 \\pm 2\\ {\\rm M}_{\\sun}$ for this star, which is within the expected mass range for SN~II-P progenitors. The yellower color could also arise from the blend of two or more stars, such as a red supergiant and a brighter, blue supergiant. Such a red supergiant hidden in this blend could instead be the progenitor and would also have an initial mass within the expected progenitor mass range. Furthermore, the yellow supergiant could be in a massive, interacting binary system, analogous to the possible yellow supergiant progenitor of the high-luminosity SN II-P 2004et. Finally, if the yellow supergiant is not the progenitor, or is not a stellar blend or binary containing the progenitor, then we constrain any undetected progenitor star to be a red supergiant with $ M_{\\rm ini} \\lesssim 11\\ {\\rm M}_{\\sun}$, considering a physically more realistic scenario of explosion at the model endpoint luminosity for a rotating star. ",
        "introduction": "Massive stars are thought to evolve to an end state which results in the collapse of the stellar core, as the hydrostatic pressure in the core can no longer support gravity.  Although the physics of what follows is still uncertain, it appears that, following core collapse, the pressure due to neutrinos released from the core results in the explosion of the rest of the star --- a supernova (SN; e.g., \\citealt{marek09}). Core-collapse supernovae (CC-SNe) have been at the forefront of astronomical research for the better part of a century, and yet we still do not have complete knowledge of which stars explode. It is possible that the relative heterogeneity in observed properties of CC-SNe arises from a range of massive progenitors.  Although indirect evidence (spectral features, light-curve shapes, environmental information) provides some clues, direct identification of the star before the SN explosion \\citep[e.g.,][]{vandyk99,vandyk00,smartt02} is the more precise means of determining the nature of the progenitor star.  At this time of writing, the only kind of CC-SNe for which progenitors have been directly identified is the Type II-Plateau supernovae (SNe~II-P; e.g., \\citealt{vandyk03,li06}), with the exception of SN~2005gl in NGC~266, a Type II-narrow supernova (SN~IIn) for which a progenitor may well have been identified \\citep{galyam07,galyam09}. The SNe~II-P exhibit Balmer lines with prominent P~Cygni-like profiles in their optical spectra at early times.  Their optical light-curve shape shows a characteristic ``plateau'' phase, prior to an exponential tail (e.g., \\citealt{barbon79}).  The trend is for their progenitor masses to be in the range of 8--20 $M_{\\odot}$ (\\citealt{li06}). Based on the statistics of 20 SNe~II-P for which progenitors have been isolated or upper mass limits established, \\citet{smartt09} derive a more limited range of $8.5 (^{+1}_{-1.5}) \\lesssim M_{\\rm ZAMS} (M_{\\odot}) \\lesssim 16.5 (\\pm 1.5)$ for these stars. All of these progenitors exploded in the red supergiant (RSG) phase, as we would theoretically expect, except possibly for SN 2004et \\citep{li05} and SN 2006ov \\citep{li07}. Given that fewer than 10 progenitors for SNe II-P have actually been directly identified, the addition of further examples is essential for validating this trend and further constraining stellar evolutionary theory.  The explosion of \\object{SN~2008cn} provides us with an opportunity to increase the sample, including particularly the high-luminosity SN subclass (e.g., \\citealt{pastorello03}), and to extend our techniques to the study of stars which have exploded in relatively more distant galaxies.  \\citet{martin08} discovered SN~2008cn on 2008 May 21.52 (UT dates are used throughout this paper), at 4$\\parcsec$7 E and 23$\\parcsec$2 N of the center of the nearby ($cz = 2592$ \\kms), nearly face-on spiral galaxy \\object{NGC~4603}.  Nine days later, \\citet{strit08} classified the SN as Type II, a few days past explosion.  \\citet{li08} measured a precise position for a putative progenitor of SN~2008cn as $\\alpha=12^{\\rm h} 40^{\\rm m} 55{\\fs}64, \\delta=-40\\arcdeg 58\\arcmin 12{\\farcs}9$ (J2000.0), from high-quality {\\sl Hubble Space Telescope\\/} ({\\sl HST}) archival images of the host galaxy taken with the Wide Field Planetary Camera~2 (WFPC2).  The position was established after the {\\sl HST\\/} images were registered with coadded $V$-band images taken with the 1.0-m Swope telescope at Las Campanas Observatory on 2008 June 1.15 (a geometrical transformation was performed, resulting in an accuracy of $0{\\farcs}1$ on the WFPC2 images). The claim was that the detected object could be a RSG with an absolute $I$-band magnitude of $-7.9$ and an intrinsic $V-I$ color of 1.2 mag.  Here we report on a more detailed analysis of this object.  In \\S \\ref{specph} we determine the characteristics of SN~2008cn. Section \\ref{identification} explains the process of identification and analysis of the possible progenitor using {\\sl HST\\/} images taken both before, and after, the SN explosion. We discuss the nature of the progenitor, based on this analysis, in \\S \\ref{discussion}, and in \\S \\ref{conclusions} we summarize our conclusions. ",
        "conclusions": "We have shown, based on a limited dataset, that the SN II-P 2008cn in NGC~4603 appears to be of the high-luminosity subclass, similar to SN~1996W \\citep{pastorello03}, SN~2003hn \\citep{krisciunas09}, and possibly SN 2004et \\citep{sahu06}.  Its $V$-band luminosity some weeks after explosion is $\\sim$0.5 mag brighter than that of, for example, the normal SN II-P~1999em \\citep{hamuy01,leonard02}. The progenitors of normal SNe II-P have all been determined or inferred to be RSGs \\citep{smartt09}.  It is unclear at this point, due to a relative lack of examples, whether high-luminosity SNe II-P, such as SN 2008cn, also arise from RSGs.  High-precision relative astrometry has been employed to identify a candidate progenitor for SN 2008cn, based on a comparison of high spatial resolution {\\sl HST\\/} imaging both of the SN at late times and of the host galaxy containing the SN site prior to explosion. Comparison of the difference between the SN and candidate positions with the astrometric uncertainties suggests that we may well have identified the SN progenitor, although this is not entirely certain.  In estimating an initial mass for the progenitor candidate, we have been careful to transform the observed photometry into intrinsic stellar quantities, specifically the bolometric luminosity and the effective temperature of the star.  We have also considered for comparison the most recent available stellar evolutionary models for massive stars, including those with rotation (in particular, the Geneva group models; the Padova group models do not include rotation). From this we have estimated an initial mass for the progenitor of SN~2008cn of $15 \\pm 2\\ {\\rm M}_{\\sun}$, which is within the initial mass range for SN II-P progenitors, based on all previous attempts that have been made to identify the star \\citep{smartt09}.  However, the star's color is more yellow than we would expect.  This unusual color could arise from the star exploding while experiencing an evolutionary blue loop, which is theoretically not expected to occur. Alternatively, the stellar profile could consist of two or more stars in a blend, with a possible combination of one or more blue and red supergiants resulting in the yellow color. In this case, a blended stellar profile could ``conceal,\" for example, a $\\sim 15\\ {\\rm M}_{\\sun}$ RSG, which could nominally be the progenitor. Additionally, the blend could be a massive, interacting binary, analogous to the yellow supergiants recently found in extragalactic eclipsing binary systems \\citep{prieto08}. That SN 2008cn, like SN 2004et, is a high-luminosity SN II-P makes this a particularly intriguing possibility. The progenitor of SN 2004et has been identified as a yellow supergiant \\citep{li05}, also coincidentally with $M_{\\rm ini} \\approx 15\\ {\\rm M}_{\\sun}$, although this identification is controversial (see \\citealt{smartt09}). Close interaction in such a massive binary system could prevent the primary from evolving to the RSG phase, allowing the star to explode as a yellow supergiant.  SN~2008cn could well represent an additional example of what is likely a relatively rare event.  Establishing the precise position of the faint progenitor candidate in the pre-explosion archival images has proven to be challenging, especially with the presence of the two brighter stars near the fainter candidate and the relatively large distance to the host galaxy. The slight excess in the positional difference in declination, relative to the positional errors, prevents us from definitively assigning this star as the progenitor. If the candidate yellow supergiant is {\\it not\\/} the SN~2008cn progenitor, then the upper limits to detecting any star within $\\sim$10 pixels ($\\sim 80$ pc) of this candidate result in a conservative upper limit to the initial mass of a RSG progenitor of $M_{\\rm ini} \\lesssim 14\\ {\\rm M}_{\\sun}$. Ultimately, the best test of the nature of the SN~2008cn progenitor is to observe this field again with {\\sl HST\\/} (or the {\\sl James Webb Space Telescope}) several years later, when the SN has faded to nondetectability. The definitive indication would be if the yellow star has vanished, or is significantly diminished in brightness, or has survived intact.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "0910/0910.0017_arXiv.txt": {
        "abstract": "Galactic weakly interacting massive particles (WIMPs) may scatter off solar nuclei to orbits gravitationally bound to the Sun. Once bound, the WIMPs continue to lose energy by repeated scatters in the Sun, eventually leading to complete entrapment in the solar interior. While the density of the bound population is highest at the center of the Sun, the only observable signature of WIMP annihilations inside the Sun is neutrinos.  It has been previously suggested that although the density of WIMPs just outside the Sun is lower than deep inside, gamma rays from WIMP annihilation just outside the surface of the Sun, in the so called WIMP halo around the Sun, may be more easily detected. We here revisit this problem using detailed Monte Carlo simulations and detailed composition and structure information about the Sun to estimate the size of the gamma-ray flux. Compared to earlier simpler estimates, we find that the gamma-ray flux from WIMP annihilations in the solar WIMP halo would be negligible; no current or planned detectors would be able to detect this flux. ",
        "introduction": "The observational evidence for dark matter is overwhelming (see e.g.\\ \\cite{Bergstrom:2000pn}), but in spite of that, direct evidence for dark matter particles has not yet been found. One of the most popular classes of dark matter models is the Weakly Interacting Massive Particle (WIMP) model. WIMPs are heavy particles with weak interaction strengths, and are popular because they are naturally thermally produced at right abundances in the early Universe (see e.g.\\ \\cite{Bergstrom:2009ib}). There are many ways to search for WIMPs, e.g.\\ via direct detection, or via indirect detections, like neutrinos from the Sun \\cite{1985ApJ...296..679P,1985PhRvL..55..257S} or the Earth \\cite{1986PhLB..167..295F,1986PhRvD..33.2079K,1986PhRvD..34.2206G}, gamma-rays from annihilations in the Galactic halo or charged cosmic-rays from annihilations in the Galactic halo (for a review, see \\cite{Bergstrom:2000pn}). The neutrino signal from the Sun arises from WIMPs scattering in the Sun, eventually sinking to the core and annihilating, producing many standard model particles. Of these, only neutrinos escape the Sun and the neutrinos can be searched for with neutrino telescopes like e.g.\\ IceCube \\cite{2004APh....20..507A}.  We will here focus on another idea, proposed by Strausz \\cite{Strausz}. The idea is that while WIMPs are being captured by the Sun, they will move on bound orbits outside the Sun, making up a halo of WIMPs around it. Pair-wise annihilations of WIMPs in this halo could occur and (if the density would be high enough), produce a detectable flux of e.g.\\ gamma-rays. This type of gamma-ray flux has been searched for by the Milagro \\cite{milagro} detector.  Strausz came to the conclusion that this gamma-ray signal would be measurable or that if no such signal is seen, one would be able to constrain parameters in the WIMP model. However, Strausz' calculation imposes many approximations and in a later preprint by Hooper \\cite{Hooper}, it was found that the flux of gamma-rays is much smaller than Strausz' estimates and a detection of this gamma-ray signal would require unrealistically large telescope areas. In both these calculations many simplifying approximations are made; for example radial orbits are assumed, the WIMPs real orbits are not followed with realistic solar models. Another calculation was performed by Fleysher \\cite{Fleysher:2003iya}, finding even higher gamma-ray fluxes than in Strausz' calculation. However, many simplifying assumptions were made in all of these calculations, making the results unreliable.  To clarify the discrepancies between these earlier estimates, we have here performed a much more thorough calculation of the WIMP density around the Sun via detailed Monte Carlo simulations. We let WIMPs from the Galactic halo scatter in the Sun and follow the WIMPs on their orbits (taking full account of the non-Keplerian nature of the orbits inside the Sun), letting them interact repeatedly with solar nuclei until they are fully trapped inside the Sun. We calculate how much each WIMP we simulate contributes to the WIMP density around the Sun. By summing up the contributions from the entire WIMP population, we can estimate the total WIMP density around the Sun. We finally calculate the gamma-ray signal from WIMP annihilations in this halo, which can be compared with the Milagro searches for a gamma-ray signal of this kind \\cite{milagro}. We will also calculate predicted fluxes for future gamma-ray detectors and compare with expected backgrounds. ",
        "conclusions": "This work is under the assumption that the WIMPs in the solar halo are not disturbed during the capture process. In reality the Solar System also has planets, whose gravitational interaction can disturb the orbiting WIMPs. This effect, as caused by Jupiter, has been discussed in e.g. \\cite{Peter}. A massive planet, like Jupiter, can disturb orbiting WIMPs so that they end up on orbits no longer intersecting the Sun, potentially making them very long lived. On the other hand, WIMPs on orbits stretching so far out spend almost all their time far away from the Sun and hence do not contribute so much to the density close to the Sun. Furthermore, planetary interactions are more likely to throw WIMPs out of the Solar System, reducing the WIMP density. WIMPs disturbed to bound orbits not intersecting the Sun end up on orbits sensitive to planetary interactions and will, in time, be disturbed again and typically lose their future contributions to the density close to the Sun. There is, however, a possibility that some fraction of these WIMPs on non-solar-crossing orbits turn out to be more stable, rendering a possible source for higher WIMP densities in the Solar System. Such a mechanism has been discussed in \\cite{dk1,dk2}, where it was argued that WIMPs which scatter in the outskirts of the Sun can be perturbed by the planets to orbits not crossing the Sun, in such a way that the new orbits become stable over very long timescales. This WIMP population would then build up over time and it was argued in \\cite{dk1,dk2} to be a significant source of WIMPs in the Solar System. These calculations were analytical, but they have not been confirmed in later numerical simulations \\cite{ap1}. In fact, in \\cite{ap1}, it is argued that the life-times would be much shorter and that the density enhancement would be quite small. In principle, it would be interesting to add planetary interactions to our simulations, as these could in principle increase the density somewhat (even if not as much as originally argued in \\cite{dk1,dk2}). That is however outside the scope of this paper, as it would be a very extensive task to do. It is also unlikely that the effects would be large enough to be interesting.  Another possibility to get higher WIMP densities is from a dark matter disk \\cite{Read:2009iv}, which could enhance the capture rate in the Sun by up to an order of magnitude \\cite{Bruch:2009rp}. This is predominantly because WIMPs from the dark disk pass at lower velocities and are hence easier for the Sun to capture as the energy loss required in the scatter is lower. However, lower initial velocities also imply that the just captured WIMPs will typically have lower energies, altering the distribution in Fig.~\\ref{fig:scatter}. This implies that the captured WIMPs will survive less scatters before being fully trapped inside the Sun, which could reduce the density contribution per captured WIMP. Hence, the dark disk scenario need not increase the density around the Sun as much as it increases the capture rate. Even under the optimistic assumption that the density is effected as much as the capture rate, which would increase the gamma-ray signal by up to  two orders of magnitude, the signal is still too low to be detectable.  The calculated gamma-ray flux at Earth from WIMP annihilations around the Sun is extremely low, as seen in table \\ref{table:flux}. The calculated flux, at best, corresponds to around one gamma photon per (10 km)$^2$ per century, which is hardly ever detectable.  The low calculated signal in this work is in contradiction with the earlier result published by Strausz \\cite{Strausz}. The analysis conducted here is, however, far more detailed. In \\cite{Strausz} it is assumed that the WIMPs move in one-dimensional orbits and that they all have some sort of average properties, such as the average energy loss in a scatter and the average scatter probability in a passage through the Sun. It is, however, very unlikely for the assumptions in \\cite{Strausz} to, by themselves, be responsible for such a strong disagreement. The passage where most of the disagreements appear is not explained in detail in \\cite{Strausz}, and hence, it is impossible to tell what the source is of the difference in results.  One should note also, that a calculation of the WIMP densities in the halo has also been performed by Fleysher \\cite{Fleysher:2003iya}. Fleysher finds even larger densities than Strausz, but that calculation cannot be correct. The solar halo WIMP densities in the calculation by Fleysher are proportional to the scattering cross sections, which is unreasonable, as the scattering cross sections that enter in the probability of the first scatter and the lifetime of the WIMP orbits cancel (see section \\ref{sec:sigmadep} for a more thorough discussion about this).  The conclusion of the work presented here, on the other hand, agrees with the conclusion argued by Hooper \\cite{Hooper}. That work is also not very detailed but more recent than Strausz' paper. Also Peter \\cite{Peter} has performed simulations of WIMPs in the Solar System, from which the WIMP density can be extracted. As discussed in \\cite{Peter:2009qj} those results are very similar to the results found here. Another test of the validity of our results is performed in Appendix \\ref{sec:reasonable}.  Finally, we can note that we set out this work with the hope of proving Strausz' optimistic calculations to be correct, but unfortunately, our detailed calculations instead show that the fluxes are indeed far too small to be detectable. On top of that, there are astrophysical backgrounds in which the signal is very well hidden. To reach this conclusion we have made optimistic assumptions about the WIMP dark matter candidate. However, we have assumed that the WIMPs are thermally produced in the early Universe, i.e.\\ has an annihilation cross section of the order of $3\\times 10^{-26}$ cm$^{3}$ s$^{-1}$. Of course, non-thermal productions or other enhancements (like a Sommerfeld enhancement at low annihilation velocities) would also be possible, but given the large enhancements needed to get an observable flux of gamma-rays from WIMP annihilations around the Sun, there are far better ways to search for those WIMPs in those cases, like gamma-rays from dwarf galaxies, annihilations in the Galactic halo, or effects on the cosmic microwave background radiation \\cite{Galli:2009zc}.  \\appendix"
    },
    "0910/0910.3157.txt": {
        "abstract": "{}{}{}{}{} % 5 {} token are mandatory  \\abstract % context heading (optional) % {} leave it empty if necessary {} % aims heading (mandatory) {We present a study of the ionized, neutral atomic, and molecular  gas associated with the ring nebula RCW\\,78 around the WR star HD\\,117688~(= WR\\,55) with the aim of analyzing the distribution of the associated gas  and  investigating its energetics. } % methods heading (mandatory) {We based our study on $^{12}$CO(1-0) and $^{12}$CO(2-1) observations of the brightest section of the nebula carried out with the SEST telescope with angular resolutions of 45\\arcsec\\ and 22\\arcsec, respectively; and on complementary $^{12}$CO(1-0) data of a larger area obtained with the NANTEN telescope with  an angular resolution of 2\\farcm 7, H{\\sc i} 21-cm line data taken from the ATCA survey, IRAS HIRES data, and radio  continuum data at 4.85 GHz from the Parkes survey. } % results heading (mandatory) {We report the detection of molecular gas having velocities in the range --56 to --33 \\kms\\ associated with the western region of RCW\\,78.  A few patches of molecular gas possibly linked to the eastern faint section are detected. The CO emission appears concentrated in a region of 23\\arcmin $\\times$18\\arcmin\\ in size, with a total molecular mass of (1.3$\\pm$0.5)$\\times$10$^5$ \\msun, mainly connected to the western section. The analysis of the neutral atomic gas distribution reveals the \\hi\\ envelope of the molecular cloud, while the radio continuum emission shows a ring-like structure, which is the radio counterpart of the optical  nebula. The gas distribution is compatible with the western section of RCW\\,78 having originated in the photodissociation and ionization of the molecular gas by the UV photons of the WN7 star HD\\,117688, and with the action of the stellar winds of the WR star on the surrounding gas. In this scenario, the interstellar bubble expanded  more easily  towards the E than towards the W due to the lack of dense molecular gas in the eastern section. The proposed scenario also explains the off center location of WR\\,55.  A number of infrared point sources classified as YSO candidates showed that stellar formation activity is present in the molecular gas linked to the nebula. The fact that the expansion of the bubble have triggered star formation in this region can not be discarded. } % conclusions heading (optional), leave it empty if necessary {} ",
        "introduction": "Interstellar bubbles created by the stellar winds of massive stars are generally detected as thermal radio continuum shells, as cavities and expanding shells in the H{\\sc i} 21-cm line emission distribution, and as  shells in the far infrared (see Cappa 2006 for a summary). Radio observations allowed to estimate the electron densities and ionized masses for a number of galactic ring nebulae and to identify their neutral gas counterparts.  The dense gas linked to these nebulae has been detected through molecular line observations.  However, up to present, only a few cases have been analyzed. These observations  revealed the existence of large amounts of molecular gas  and that photodissociation regions at the interface between the ionized and molecular gas and shock fronts are present in interstellar bubbles (Cappa et al. 2001, Rizzo et al. 2003).  In the last years,  infrared images obtained from the MSX Galactic Plane Survey (Price et al. 2001) and GLIMPSE survey (IRAC images, Benjamin et al. 2003) allowed to investigate the  distribution of the near- and mid-infrared emission associated with interstellar bubbles, confirming that PDRs  and shock fronts are common phenomena linked to these structures (Churchwell et al. 2006,  Cyganowski et al. 2008, Watson et al. 2008, Watson et al. 2009).  Observational and theoretical studies, as well as numerical simulations, showed that only a small amount of the stellar wind energy released to the interstellar medium is converted into kinetic energy of the bubbles (see for example Chu et al. 1983, Oey 1996, Freyer et al. 2003, 2006, Cooper et al. 2004, Cappa 2006).  An important issue is the fact that the stellar formation process may be favoured in the compressed layers around the nebulae (Elmegreen 2000, Thompson et al. 2004).  Using infrared point source catalogues, signs of stellar formation activity have been found in the dense molecular envelopes surrounding some interstellar bubbles (e.g. Cappa et al. 2005, Zavagno et al. 2007).  Thus, observations of molecular gas associated with interstellar bubbles provide information essential to investigate  the energetics and stellar formation process in the dense  surrounding shells.  Here we investigate the distribution of the molecular gas associated with the optical ring nebula RCW\\,78 around WR\\,55 based on high angular resolution SEST observations, complemented with lower resolution data of the NANTEN telescope. Complementary IR, \\hi\\ 21-cm line, and radio continuum archival data allow the analysis of the distribution of the ionized and neutral atomic gas, and that of the dust. Our aims are to identify and characterize the material linked to the ring nebula and to study its kinematics and energetics.  RCW\\,78 (= G\\,48b) is a ring nebula of about 35\\arcmin\\ in diameter. The DSS-R image  of the whole nebula is shown in Fig. 1, where the white cross marks the position of the WR star (see also figure 5 by Chu et al. 1983). The brightest part of RCW\\,78 is about 10\\arcmin \\por 6\\arcmin\\ in size and offset to the northwest of the star, while fainter regions are present to the northeast, east, and south (Chu \\& Treffers 1981, see also Heckathorn et al. 1982). Chu (1981) classified the optical nebula as R$_a$ because of its low expansion velocity and the lack of evidence for a shell structure in the velocity pattern of the bright regions. This last characteristic is explained by Chu \\& Treffers as due to the fact that the stellar winds have existed during a short period of time.  %-----------------------------------------------figure 1------------ \\begin{figure} \\centering \\includegraphics[width=8.5cm]{ccappa.fig1-b.eps} \\caption{DSS-R image of the eastern and western sections of RCW\\,78. The cross indicates the location of the WR star. } \\label{FigVibStab} \\end{figure} %%---------------------------------------------figure 1------------  The H$\\alpha$ study by Chu \\& Treffers (1981) showed that the velocity of the ionized  material spans from --38 to --53 \\kms, varying from --44 \\kms\\ near the star to --53 \\kms\\ 7\\arcmin\\ north of it. Georgelin et al. (1988) found a similar velocity of --41.4 \\kms\\ based on H$\\alpha$ Fabry-Perot data. Chu \\& Treffers (1981) explained the observed velocity pattern, which does not correspond to that of an expanding shell, as the consequence of an outflow caused by the ionization of the surface of the molecular cloud by the central WR star. Georgelin et al. (1988) found an additional component in the H$\\alpha$ line profiles at --24 \\kms,  most probably unrelated to the nebula.   %-----------------------------------------------figure 2------------ \\begin{figure*} \\centering \\includegraphics[width=15cm]{ccappa.fig2-b.eps} \\caption{The central panel displays the SuperCOSMOS H$\\alpha$ image of the W section of RCW\\,78. The grayscale is in arbitrary units. The crosses indicate the position of the 170 points observed in CO lines. The $^{12}$CO(2-1) spectra corresponding to selected positions show molecular components with different velocities. The position of each spectrum is indicated. Intensities are given in main-beam brightness temperature.} \\label{FigVibStab} \\end{figure*} %%---------------------------------------------figure 2------------  Circular galactic rotation models (e.g. Brand \\& Blitz 1993) predict that gas with LSR velocities in the range --38 to --53 \\kms\\ is placed at kinematical distances of 3.5-7.0 kpc. In this section of the Galaxy, the mentioned velocities are close to that of the tangent point, i.e. --48 \\kms.  The study of the ionization structure by Esteban (1993) indicates that photoionization is the main source of excitation of the nebula, compatible with the clasification by Chu (1981).  The nebula is related to HD\\,117688 (= WR\\,55 = MR\\,49),  a WN7 star located at {\\itshape (l,b)} = (307\\deg 48\\arcmin,+0\\deg 9\\farcm 6) or \\radec\\ = (13$^h$33$^m$30.1$^s$, --62\\deg 19\\arcmin 1.2\\arcsec). Spectrophotometric distances $d$ were estimated by several authors: 4.0 kpc (Georgelin et al. 1988), 5.5 kpc (Conti \\& Vacca 1990), 6.0 kpc (van der Hucht 2001). New NIR calibrations of absolute magnitudes in the  $K_s$ band by Crowther et al. (2006), indicate an absolute magnitude $M_{Ks}$ = --5.92 mag for WN7-9 weak lined stars. Taking into account the  $K_s$-value for WR\\,55 from the 2MASS catalogue (Cutri et al. 2003), and interstellar extinction values  from Marshall et al. (2006), a distance in the range 4.5-5.0 kpc can be derived. Based on the available distance estimates, we will adopt a distance of 5.0$\\pm$1.0 kpc for the WR star and its surrounding ring nebula.  The terminal wind velocity for WR\\,55 derived from the P Cyg profile of the CIV $\\lambda$1550 line is in the range $V_w$ = 1000 -- 1200 \\kms\\ (Hamann et al. 1993, Rochowicz \\& Niedzielski 1995, Niedzielski \\& Sk\\'orzy\\'nski 2002). As regards the mass loss rate, Hamann et al. (1993) estimate $log$ \\.M  = --4.2 \\msunyr.  In this paper, we present high angular resolution SEST observations performed in the $^{12}$CO(1-0) and $^{12}$CO(2-1)  lines towards the brightest section of the nebula, complemented with lower angular resolution NANTEN data, \\hi\\ 21-cm line data, radio continuum images, and infrared data of the whole region. In the next sections we describe the SEST data and analyze the distribution of the molecular, ionized, and neutral atomic emissions.  %__________________________________________________________________ ",
        "conclusions": "We have investigated the distribution of the molecular gas related to the ring nebula RCW\\,78 around the WR star HD\\,117688 based on observations of $^{12}$CO(1-0) and $^{12}$CO(2-1) lines obtained using the SEST. Complementary molecular data from the NANTEN survey, \\hi\\ data from the SGPS, and IRAS and radio continuum images were also analyzed.  The analysis of the SEST and NANTEN CO data reveals the existence of molecular gas having velocities in the range --56 to --33 \\kms\\ interacting with the western section of the nebula. The bulk of the emission is concentrated in two structures having velocities in the range  --52.5 to --43.5 \\kms\\ and --43.5 to --39.5 \\kms, coincident with the velocities of the H$\\alpha$ line. Based on morphological and kinematical evidences we conclude that the arc-like structure having velocities in the range --52.5 to --43.5 \\kms\\ is clearly associated with RCW\\,78. The feature with velocities from --43.5 to --39.5 \\kms\\ is most probably related to the nebula. On the contrary, CO emission linked to the  eastern section of the nebula is difficult to identify.  The radio continuum emission distribution at 4.85 GHz shows a shell-like feature, which  resembles the H$\\alpha$ emission distribution.  The \\hi\\ emission distribution shows a low emission region coincident in position with molecular gas with similar velocities, suggesting that most of the neutral gas in this region is H$_2$.  In this scenario, part of the neutral atomic gas encircling the depression might correspond to the envelope of the molecular cloud.  The distribution of the molecular and ionized gas is compatible with a scenario in which the ionized gas originated through photodissociation and ionization of the parental molecular cloud. The shell-like appearance of the ionized gas is suggestive of the action of the stellar winds,  which swept-up the surrounding gas shaping an interstellar bubble. The lack of dense gas towards the east favored the expansion of the nebula in this direction. The proposed scenario also explains the off center location of WR\\,55.  As regards stellar formation activity towards RCW\\,78, the analysis of the IRAS, MSX, and Spitzer point source catalogues revealed the existence of two active star forming regions liked to the molecular gas associated to the nebula. The influence of the expansion of RCW\\,78 in the onset of star formation in these regions can not be discarded.  The comparison between the wind mechanical energy released by WR\\,55 and its massive progenitor and the kinetic energy of the interstellar bubble indicates that the stellar wind of this star is capable of blowing the interstellar bubble. Thus, WR\\,55 is not only responsible for the ionization of the gas in the nebula but for the creation of the interstellar bubble."
    },
    "0910/0910.4825_arXiv.txt": {
        "abstract": "We study cosmological constraints on the various accelerating models of the universe using the time evolution of the cosmological redshift of distant sources. The important characteristic of this test is that it directly probes the expansion history of the universe. In this work we analyze the various models of the universe which can explain the late time acceleration, within the framework of General Theory of Relativity (GR) (XCDM, scalar field potentials) and beyond GR ($f(R)$ gravity model). ",
        "introduction": "The recent accelerated expansion of the universe is one of the most important discovery in the cosmology. Whether this observed acceleration is due to some new hypothetical energy component with large negative pressure (dark energy) within the framework of General  Theory of Relativity, or due to modification in the GR at the cosmological distances (modified gravity), is not known. Therefore many models have been proposed in literature to understand the origin and nature of this present acceleration~\\cite{da}.  Basically the cosmic acceleration affects the expansion history of the universe. Therefore to understand the true nature of this driving force, mapping of the cosmic expansion of the universe is very crucial~\\cite{li}. Hence, we require various observational probes in different redshift ranges to understand the expansion history of the universe. The observational tools for probing the cosmic acceleration broadly fall into two categories: Geometrical and Dynamical probes. A {\\it Geometrical probe} deals with large scale distances and volume which include luminosity distance measurements of SNe Ia, angular diameter distance from first CMB acoustic peak, Baryon Acoustic Oscillations (BAO) etc.. A {\\it Dynamical probe} investigates the growth of matter density perturbations that give rise to the  large scale structure such as galaxies, clusters of galaxies etc. in the universe.  Using supernovae as standard candles is a popular method of constraining the properties of dark energy. Though this method is very simple and useful in constraining the various dark energy models, at present the luminosity distance measurements suffers from many systematical uncertainties like extinction by dust, gravitational lensing etc.~\\cite{nordin}. On the other hand, measuring the expansion history from growth of matter perturbations also has its limitations. It requires prior information of exact value of matter density, initial conditions, cosmological model etc.~\\cite{li,l}. So the question arises, ``Is there any probe which is simple, depends on fewer priors and assumptions?'' The possible probe is ``Cosmological Redshift Drift'' (CRD) test which maps the expansion history of the universe directly.  The CRD test is based on very simple and straightforward physics. However, observationally it is a very challenging task and requires technological breakthroughs~\\cite{odorico}. The most remarkable feature of this probe is that it measures the dynamics of the universe directly - the Hubble expansion factor. This property makes it very special and unique. Further, it assumes that the universe is homogeneous and isotropic at the cosmological scales [for more details see ref.~\\cite{lis}]. The time drift of the cosmological redshift probes the universe in the redshift $2 < z < 5$, whereas the other cosmological tests based on SNe Ia, BAO, weak lensing, number counts of clusters etc. have not penetrated beyond $z = 2$. The other advantage of this tool is that it has controlled systematical uncertainties and evolutionary effects of the sources.  The aim of this letter is to employ CRD test to constrain various accelerating models both within the framework of GR and beyond GR. We have used \\emph{simulated data points} for redshift drift experiment generated by Monte Carlo simulations with the assumption of standard cosmological model ($\\Lambda$ CDM) as reference~\\cite{lis,joe}. We put constraints on the Quintessence models (based on scalar field potentials) like PNGB, inverse  power law and exponential potentials. The dark energy parametrization which has a variable equation of state is also investigated. We also put a bound on the $f(R)$ gravity models i.e., Starobinsky model using the Cosmological Redshift Drift test.  The letter is organized as follows. In Section 2 we review the theoretical basis of Cosmological Redshift Drift test. We also present here methodology and data used for this work. Various models which can explain late time accelerated expansion of the universe are described in Section 3. Last section contains a summary and discussion of results. ",
        "conclusions": "Till now, large number of theoretical models have been proposed to explain the observed accelerated expansion of the universe. In order to get insight of the mechanism behind this acceleration, we require complimentary observational tools. The observational tests either belong to {\\em distance based} methods such as SNe Ia luminosity distances, angular size of compact radio sources, BAO, CMBR, gravitational lensing  etc. or {\\em time based} methods like Absolute age method, Lookback time method and  Differential age method. Every observational tests  requires some priors, assumptions and is subjected to systematic errors.   There is a need to develop tools which are simple and have a controlled systematics. The cosmological redshift drift (CRD) method is a test in this direction. Corasaniti et al. were the first to analyse dark energy models like $\\Lambda$CDM, Chaplygin gas and dark energy- dark matter interaction model using the CRD test~\\cite{cor}. They conclude that the CRD test puts stringent constraints on non-standard dark energy models. Later on, Balbi and Quercellini also investigated various standard and non-standard dark energy models~\\cite{ba}. They found  the worst bound for the Cardassian model.  Further Zhang et al. also studied Holographic dark energy model with CRD test~\\cite{zh} and obtain a very tight bound on the $\\Omega_m$.  The above stated work used the data generated by MC simulations, given by Pasquini et al. (2006)~\\cite{pa}.  Uzan, Bernardeau and Mellier discuss the possibility in which the large scale structure may effect the measurements of  time drift of the cosmological redshift~\\cite{u}. Since the CRD test is based on the assumption that the universe is homogeneous and isotropic, therefore this test has been used to check the homogeneity of the universe~\\cite{u1}. Quartin and Amendola propose that CRD test can be used to distinguish between Void models and conventional dark energy scenarios~\\cite{q}. Recently one possible source of noise in the CRD test has been studied by Killedar and  Lewis which deal with the transverse motion of the Ly$\\alpha$ absorbers \\cite{kill}.  Recently Liske et al. (2008) studied in detail the impact of next generation ELT on observing the very small redshift drift signal~\\cite{lis}. Using the extensive MC simulations, three sets of data points are generated with  different strategies which include the measurement of the precise value of velocity shift. In  our work, We used the recent data generated by Liske at al.(2008) to constrain the late time acceleration models of the universe. We study the models which belong to both the Einstein gravity scenario ($\\Lambda$CDM, XCDM, scalar field potentials) and the modified gravity ($f(R)$).  Results are summarized as follows:  \\begin{enumerate}  \\item In $\\Lambda$CDM model $(w_x = -1)$, the $\\chi^2$ minimum lies at $\\Omega_{m0}= 0.3$. Considering $h$ to be a nuisance parameter, we marginalize over $h$ to obtain the probability distribution function defined as:  $$ L(p)\\, = \\, \\int{e^{-\\chi^2(h,p)/2}P(h)dh} $$  Here P(h) is the prior probability function for $h$ which is assumed to be Gaussian. We find that the best fit value of model  parameter is independent of the choice of prior.  \\item We performed the $\\chi^2$ statistics on flat XCDM dark energy model as shown in Fig. \\ref{contour}. The best fit values lies at  $\\Omega_{m} = 0.3$ and  $ w_x = - 1$. At $1 \\sigma $-level the constraints are, \\[ \\Omega_m \\, = \\, 0.30^{+0.03}_{-0.04} \\hskip 2 cm \\omega_x \\le -0.62 \\] Here again we marginalize over $h$ to obtain the best fit value of model  parameters $(\\Omega_m, w_x)$. The best fit values are independent of the choice of prior. The redshift drift test give very tight constraint on the $\\Omega_m$ and weak bound on the equation of state, $\\omega_x$. This is expected since the amplitude and slope of the cosmic velocity shift w.r.t the redshift is very sensitive to the $\\Omega_m$ and shows weak dependence on the equation of state. The other important feature of this analysis is that it is  complementary to other probes such as CMB, BAO and weak lensing.  \\begin{figure}[ht] \\centering \\includegraphics[width=6.5cm]{contour.eps} \\caption{Flat dark energy model with constant $\\omega_x$: Contours in $\\Omega_m$ and $\\omega_x$ plane for CRD test. The best fit value lies at $ \\Omega_m= 0.3$ and  $\\omega_x = -1$. The inner and outer curves are at $1 \\sigma$ and $2 \\sigma$ respectively. }\\label{contour} \\end{figure}   \\item We analyze the $f(R)$ gravity model proposed by Starobinsky  with the CRD test. The variation of velocity drift for different values of $n$ with  redshift $z$ is shown in Fig. \\ref{fig:fR}. The expansion history of this model match exactly with the $\\Lambda$CDM model for $ n = 2$ and $\\lambda = 2$. In this model the Hubble parameter, $H(z)$, becomes independent of the parameter $n$ after the redshift $z = 1.8$ and expansion history traces exactly the $\\Lambda$CDM model behavior at high redshift~\\cite{odintsov,dev}. Since the simulated data points have $z \\geq 1.9$, the $\\chi^2$ for this model will be the same as for the $\\Lambda$CDM model. Here $\\Lambda$CDM model means standard cosmological model with $\\omega_x = -1$, $\\Omega_{m0}= 0.3$ and $\\Omega_{\\Lambda 0}= 0.7$ . The effective equation of state for this model appproaches $ \\omega = -1$ at the present epoch [see Fig.8 of the ref.\\cite{dev}]. In order to constrain this model better we need redshift drift data in the redshift range $ z \\le 1.9$ .  \\item In Fig. \\ref{chi_exp}, we plot the variation of $\\chi^2$ with the parameter $\\beta$ for the exponential potential. The best fit value of $ \\beta = m_{pl}^2/(8\\pi f^2)=3.2$ corresponds to $ \\Omega_{m0} = 0.26$. The $\\chi^2$ per degree of freedom is 0.67.  At $3 \\sigma$ limit, the allowed range of $ \\beta = 3.2^{-0.45}_ {+0.55}$ gives $ \\Omega_m = 0.26^ {-0.05}_{+0.06}$. Although the model predicts the observed value of $\\Omega_m$ but still this model is not fully in concordance with the observations; as  the allowed range of $\\beta$ at $3 \\sigma$ level shows that the universe mostly remains in a decelerating phase (see Fig. \\ref{fig:exp-pot}).   \\begin{figure}[ht] \\centering \\includegraphics[width=5.2 cm,angle=-90]{chi_expo.eps} \\caption{Variation of $\\chi^2$ with $\\beta$}\\label{chi_exp} \\end{figure}  \\item  In the inverse power law potential model, for $\\gamma \\rightarrow 0$, the energy-momentum tensor approaches towards cosmological constant. The best fit values for this model are  $\\gamma =0.025$ and $K = 0.118 $. In Fig. \\ref{Contour_Power_law} the contour is plotted in $ \\gamma - K$ plane. We would like to note here that computational efforts to generate the data points in the neighborhood of $\\gamma = 0$ increases. The important feature of this contour is that it is complementary to the contour obtain by using the Dark Energy Task Force (DETF) simulated data sets for future experiments. This includes SNe Ia, BAO, weak Lensing and CMB (PLANCK) observations \\cite{y}.  In the Fig. \\ref{eqn_power_law}  variation of instantaneous equation of state is plotted w.r.t. redshift for the best fit values of the model parameters. There is almost no variation in the equation of state with redshift. It always stay close to $ \\omega = -1$  for the redshift range studied here.   \\begin{figure}[ht] \\centering \\includegraphics[width=6.0cm,angle=0]{contour_power_law.eps} \\caption{Contour in  $\\gamma - K$ plane for inverse  power law potential.  The inner and the outer contours are  drawn at the  1$\\sigma$ and 2$\\sigma$ level respectively.}\\label{Contour_Power_law} \\end{figure}  \\begin{figure}[ht] \\centering \\includegraphics[width=6.0cm,angle=0]{es_power_law.eps} \\caption{Variation of $\\omega$ w.r.t. redshift for inverse  power law potential.}\\label{eqn_power_law} \\end{figure}   \\item For the PNGB model, the contour is plotted in $\\alpha - m$ plane as shown in the Fig. \\ref{Contour_PNGB}. The $\\chi^2$ minimum lies at $\\alpha =0.21 $ for $ m = 0.56$. At 2$\\sigma$ level, $\\alpha > 0.15$ is allowed. This result is completely  in concordance with the other cosmological observations (eg. SNe Ia, lensing, cluster observations) \\cite{waga,ng,wa}. Again this contour is complementary to the contour obtain by using Markov  Chain Monte Carlo analysis (MCMC) to generate the future data set. This includes future SNe Ia, BAO, weak Lensing and cosmic microwave background observations \\cite{aa}.  In the Fig. \\ref{eqn_pngb_law} instantaneous equation of state  is plotted for the best fit values. In contrast to power law potential, the PNGB show large variation in equation of state for the redshift $ z < 2$.  \\begin{figure}[ht] \\centering \\includegraphics[width=5.30 cm,angle=0]{contour_pngb.eps} \\caption{Contour in  $\\alpha - m$ plane for PNGB potential. The inner and the outer contours are  drawn at the  1$\\sigma$ and 2$\\sigma$ level respectively.}\\label{Contour_PNGB} \\end{figure}  \\begin{figure}[ht] \\centering \\includegraphics[width=5.30cm,angle=0]{es_pngb.eps} \\caption{Variation of  $\\omega $ w.r.t. redshift for for PNGB potential.}\\label{eqn_pngb_law} \\end{figure}  It is clear that CRD test is a very simple, straightforward and powerful tool which probes the expansion history of the universe directly. In future the precision data especially in redshift range of $z < 2$ shall add to the predictive power of CRD test.  \\end{enumerate}"
    },
    "0910/0910.1588_arXiv.txt": {
        "abstract": "We investigate indirect neutrino signals from annihilations of Kaluza--Klein dark matter in the Sun. Especially, we examine a five- as well as a six-dimensional model, and allow for the possibility that boundary localized terms could affect the spectrum to give different lightest Kaluza--Klein particles, which could constitute the dark matter. The dark matter candidates that are interesting for the purpose of indirect detection of neutrinos are the first Kaluza--Klein mode of the $U(1)$ gauge boson and the neutral component of the $SU(2)$ gauge bosons. Using the DarkSUSY and WimpSim packages, we calculate muon fluxes at an Earth-based neutrino telescope, such as IceCube. For the five-dimensional model, the results that we obtained agree reasonably well with the results that have previously been presented in the literature, whereas for the six-dimensional model, we find that, at tree-level, the results are the same as for the five-dimensional model. Finally, if the first Kaluza--Klein mode of the $U(1)$ gauge boson constitutes the dark matter, IceCube can constrain the parameter space. However, in the case that the neutral component of the $SU(2)$ gauge bosons is the LKP, the signal is too weak to be observed. ",
        "introduction": "Dark matter (DM) is one of the most active research areas in the interdisciplinary context of particle physics, astroparticle physics, and cosmology today. In recent years, experiments and observations have resulted in compelling evidence for the existence of non-baryonic DM. A good fit to experimental data is obtained if it is assumed that the DM is cold, {\\it i.e.}, non-relativistic at the epoch of matter-radiation equality, and that there is a cosmological constant accounting for the observed present expansion of the Universe. An analysis combining the five-year WMAP data, baryon acoustic oscillations, and supernova observations indicates that the DM has a density of $\\Omega_{\\rm DM} h^2 = 0.1131 \\pm 0.0034$ (68 \\% CL) \\cite{Komatsu:2008hk}. To date, all evidence for DM comes from gravitational interactions, and its particle nature remains unknown. In particular, none of the particles in the current Standard Model (SM) of particle physics could constitute the DM\\footnote{Although neutrinos only interact weakly, they are too light to be able to make up more than a small fraction of the DM density.}. Hence, the search for a DM candidate is strongly connected to physics beyond the SM.  An important class of DM candidates constitutes Weakly Interacting Massive Particles (WIMPs), which are weakly interacting particles with masses in the range $\\mathcal{O}(1 - 1000) \\, {\\rm GeV}$. The reason for the importance of WIMPs as DM candidates is that if they are assumed to be thermally produced in the early Universe, the measured present relic DM density can be easily obtained without any further strong assumptions on their properties.  Several currently operating or planned experiments aim to detect WIMPs. This could be done either directly by measuring the recoil of atoms from their elastic scattering on WIMPs \\cite{Goodman:1984dc} or indirectly by detecting products of pair annihilations or decays of WIMPs, such as gamma rays, neutrinos, and antimatter.  Through their interactions with nuclei, WIMPs in the solar system scatter elastically in the Sun, thereby losing energy and becoming trapped in the solar gravitational potential. Therefore, the local DM density in the Sun can exceed the average density by several orders of magnitude, and hence, the Sun could be an interesting source of WIMP annihilation products. In this context, only neutrinos are of interest, since they have a sufficiently low interaction strength to be able to escape the dense medium of the Sun. These neutrinos can be searched for in neutrino telescopes on Earth, such as the AMANDA/IceCube \\cite{Abbasi:2009uz,Braun:2009fr,Danninger:2009uf}, ANTARES \\cite{Lim:2009jy}, and Super-Kamiokande \\cite{Desai:2004pq} experiments.  An interesting type of WIMPs is Kaluza--Klein Dark Matter (KKDM) \\cite{Servant:2002aq,Cheng:2002ej}, which arises in certain models of extra dimensions. In particular, Universal Extra Dimensions (UED) models \\cite{Appelquist:2000nn}, where all the particles of the SM are allowed to propagate in one or more extra dimensions, could include several different candidates. In these models, a remnant of translational invariance in the extra dimensions gives rise to a multiplicatively conserved quantum number known as Kaluza--Klein parity, forcing the total KK number entering any vertex to be even. This symmetry is similar to $R$-parity in supersymmetric models, and it ensures that the lightest Kaluza--Klein particle (LKP) is stable. If neutral, the LKP could be a good DM candidate, analogously to the lightest supersymmetric particle (LSP) \\cite{Martin:1997ns}.  In the UED models, the mass spectrum of the KK particles, and in particular, the identity of the LKP, depends on specific assumptions on the model at the cutoff scale, where the theory breaks down. In the most widely studied five-dimensional model, known as the Minimal Universal Extra Dimensions (MUED) model, the LKP has been determined to be the first KK mode of the $U(1)$ gauge boson by calculating radiative corrections to the tree-level masses of the KK excitations \\cite{Cheng:2002iz}. However, as has explicitly been shown in Ref.~\\cite{Flacke:2008ne}, other assumptions could change this conclusion.  The most stringent constraints to date set on the five-dimensional MUED model of KKDM are given by the IceCube collaboration \\cite{Abbasi:2009vg}, using the 22-string configuration. When completed with in total 80 strings, IceCube will have the ability to set even stronger constraints on the model.  In this paper, we calculate muon fluxes induced by neutrinos coming from annihilation of KKDM in the Sun. We study different LKP candidates in a five- as well as a six-dimensional model. We use the DarkSUSY \\cite{Gondolo:2004sc} and WimpSim \\cite{Blennow:2007tw} packages to properly treat the capture rates, the interactions of the annihilation products in the Sun, and the propagation of neutrinos to the Earth, and to simulate the interactions of neutrinos in the detector medium on Earth.  Neutrinos from KKDM annihilations in the Sun have previously been studied in Ref.~\\cite{Hooper:2002gs}, and recently in Ref.~\\cite{Flacke:2009eu}. In comparison to those works, we use a more careful treatment, and we find relatively small differences when comparing the results. To our knowledge, the only study of neutrinos from KKDM annihilations in the Sun in a six-dimensional model has been performed in Ref.~\\cite{Dobrescu:2007ec}. In that work, only the MUED model is investigated, and it is concluded that no observable neutrino telescope signal is obtained.  The rest of the paper is organized as follows: In Sec.~\\ref{sec:WimpAnn}, we describe the formalism of WIMP annihilations in the Sun, and how this is implemented in DarkSUSY and WimpSim. In Sec.~\\ref{sec:Models}, we introduce the extra-dimensional models that we investigate. In Sec.~\\ref{sec:Results}, we present our results in the form of muon fluxes in a neutrino telescope on Earth and compare these to the existing results in the literature. Finally, in Sec.~\\ref{sec:Summary}, we summarize and discuss our results. In addition, technical details are provided in the appendices. ",
        "conclusions": "\\label{sec:Summary}  In this work, we have studied neutrinos from annihilations of KKDM in the Sun. Using the DarkSUSY and WimpSim packages, we have calculated the flux of neutrino-induced muons and antimuons in an Earth-based neutrino telescope, such as IceCube. We have investigated KKDM in a five-dimensional model, based on the orbifold $S^1/{\\mathbb Z}_2$, and in a six-dimensional model, based on the chiral square $T^2/{\\mathbb Z}_4$. Rather than restricting ourselves to the MUED models, where the LKP is uniquely determined, we have allowed for the possibility that BLTs might affect the spectrum in such a way as to change the identity of the LKP. The possible DM candidates are then the first KK modes of the neutrinos, the neutral components of the Higgs field, and the $B$ and $W^3$ gauge bosons. In addition, in the six-dimensional model, the adjoint scalars $B_H^{(1)}$ and $W_H^{3(1)}$ are possible candidates. Among these particles, neutrinos are already ruled out by direct detection experiments, while the interactions of scalars with nucleons are too strongly constrained to allow for an observable signature in neutrino telescopes. In five as well as six dimensions, only the $B^{(1)}$ and the $W^{3(1)}$ are left as interesting DM candidates.  For a given LKP mass, the flux is somewhat lower for the case of the $W^{3(1)}$ as the LKP than for the case of the $B^{(1)}$. However, it is important to note that a $B^{(1)}$ LKP is not expected to have the same mass as a $W^{3(1)}$ LKP. Relic density calculations show that the mass of the $B^{(1)}$ should lie in the range from about $500 \\, {\\rm GeV} - 1600 \\, {\\rm GeV}$, whereas the $W^{3(1)}$ should be far more massive, in the range from about $1800 \\, {\\rm GeV} - 2700 \\, {\\rm GeV}$. Since the capture rate falls off approximately as $m_{\\rm LKP}^{-6}$, this implies that the fluxes from the $W^{3(1)}$ are expected to be much smaller. In the relevant range, the flux is predicted to be smaller than $0.2 \\, {\\rm km}^{-2} \\, {\\rm yr}^{-1}$. This means that, even assuming a perfect $1 \\, {\\rm km}^2$ detector, no observable muon fluxes will be generated in IceCube. We wish to emphasize this fact, which has previously not been discussed in the literature. In contrast, in the case that the $B^{(1)}$ is the LKP, IceCube will be able to put constraints on the relevant parts of the parameter space.  Another novel conclusion is that, since the fluxes are equal for the five- and the six-dimensional model for each LKP, it is not possible to distinguish these two models using this indirect detection method.  In conclusion, in the case that the $B^{(1)}$ is the LKP, neutrino telescopes such as IceCube have the ability to set constraints on both the five- and the six-dimensional model that we have studied. On the other hand, if $W^{3(1)}$ is the LKP, the prospects for indirect detection through neutrinos are much worse. Note also, that in this case, the scattering cross section is far below the sensitivity of any currently running or planned direct detection experiments."
    },
    "0910/0910.4152_arXiv.txt": {
        "abstract": "Supermassive black hole binaries (SMBHBs) are products of galaxy mergers, and are important in testing $\\Lambda$ cold dark matter cosmology and locating gravitational-wave-radiation sources. A unique electromagnetic signature of SMBHBs in galactic nuclei is essential in identifying the binaries in observations from the IR band through optical to X-ray. Recently, the flares in optical, UV, and X-ray caused by supermassive black holes (SMBHs) tidally disrupting nearby stars have been successfully used to observationally probe single SMBHs in normal galaxies. In this Letter, we investigate the accretion of the gaseous debris of a tidally disrupted star by a SMBHB. Using both stability analysis of three-body systems and numerical scattering experiments, we show that the accretion of stellar debris gas, which initially decays with time $\\propto t^{-5/3}$, would stop at a time $T_{\\rm tr} \\simeq \\eta T_{\\rm b}$. Here, $\\eta \\sim0.25$ and $T_{\\rm b}$ is the orbital period of the SMBHB.  After a period of interruption, the accretion recurs discretely at time $T_{\\rm r} \\simeq \\xi T_b$, where $\\xi \\sim 1$. Both $\\eta$ and $\\xi$ sensitively depend on the orbital parameters of the tidally disrupted star at the tidal radius and the orbit eccentricity of SMBHB. The interrupted accretion of the stellar debris gas gives rise to an interrupted tidal flare, which could be used to identify SMBHBs in non-active galaxies in the upcoming transient surveys. ",
        "introduction": "\\label{introduction}  Supermassive black hole binaries (SMBHBs) are predicted by the hierarchical galaxy formation model in $\\Lambda$ cold dark matter ($\\Lambda$CDM) cosmology \\citep{beg80,vol03}. After a SMBHB at the center of merging systems become hard, it would stall at the hard radius for a timescale even longer than the Hubble time if spherical two-body relaxation dominates \\citep{beg80}. However, recent investigations suggested that the hardening rates of SMBHBs can be boosted and SMBHBs may coalesce within a Hubble time either due to various stellar dynamical processes other than the spherical two-body relaxation \\citep{yu02,cha03,mer04,ber06,ses08}, or due to gas dynamics \\citep[][and references therein]{gou00,liu03,cop09}.  The strong gravitational wave (GW) radiation generated by coalescing supermassive black holes (SMBHs) is the main target of the GW detector the Laser Interferometer Space Antenna (LISA) and of the Pulsar Timing Array (PTA) program. Because of the poor accuracy of both LISA and PTA in locating GW radiation sources, it is of key importance to detect electromagnetic counterparts (EMCs) of GW radiation sources. Identifying SMBHBs by their EMCs is also essential to constraining the poorly understood galaxy-merger history. Several EMCs have been suggested in the literature to probe SMBHBs and their coalescence: (1) precession of jet orientation and its acceleration in radio galaxies during the in-spiraling of SMBHBs \\citep{beg80,liu07}, (2) optical periodic outbursts in blazars due to the interaction between SMBHB and accretion disk \\citep{sil88,liu95,liu06,liu02,val08,hai09}, (3) jet reorientation in X-shaped radio galaxies due to the exchange of angular momentum between SMBHB and accretion disk \\citep{liu04}, (4) two systems of broad emission lines (BELs) in quasars \\citep{bor09}, (5) intermittent activity in double-double radio galaxies at binary coalescence \\citep{liu03}, (6) X-ray, UV, optical, and IR afterglow following binary coalescence \\citep{mil05,shi08a,lip08,sch08a}, and (7) systematically shifted BELs relative to narrow emission lines \\citep{mer06,kom08} and off-center active galactic nuclei \\citep[AGNs;][]{mad04,loe07} because of SMBH GW radiation recoil.  All the above observational signatures require gas accretion disks around SMBHBs. The SMBHBs in gas-poor galactic nuclei are difficult to detect, because the SMBHs are dormant. However, a dormant SMBH could be temporarily activated by tidally disrupting a star passing by and accreting the disrupted stellar debris \\citep{hil75}. A tidal flare decays typically as a power law $t^{-5/3}$ \\citep{ree88,phi89}, which has been observed in several non-active galaxies \\citep{kom99,kom02,hal04,esq08,gez08,gez09}. Recently, \\citet{che08} and \\citet{che09} calculated the tidal disruption rate in SMBHB systems at different evolutionary stages, and found that it is significantly different from the typical rate for single SMBHs by orders of magnitude. This great difference of the flaring rates enables one to statistically constrain the SMBHB population in normal galaxies \\citep{che08}, but identifying SMBHBs individually is still difficult at present. In this Letter, we investigate the influence of a SMBHB on the accretion of tidally disrupted stellar plasma and on the tidal flare. We show that the accretion of the stellar debris is interrupted by the SMBHB and this interruption can be taken as a key distinguishable observational signature for SMBHBs in gas-poor galactic nuclei. ",
        "conclusions": "\\label{discussion}  We investigated the accretion of tidally disrupted stellar debris in SMBHB systems. For simplicity, we assumed for the tidal debris gas: (1) a constant distribution of mass in binding energy at the tidal radius, (2) ballistic motion and negligible interaction of the fluid elements after tidal disruption, and (3) instant circularization of the debris gas probably due to shocks because of interaction between the fluid elements returned to within a radius two times the tidal radius. The first two assumptions have been justified by numerical hydrodynamic simulations \\citep[e.g.][]{eva89,lod08}, but the third one needs to be verified by hydrodynamic simulations capable of capturing strong shocks. With these three assumptions, the fluid elements and SMBHB compose restricted three-body systems, so we can investigate the accretion of the gaseous stellar debris by analyzing the stability of the three-body systems using the {\\it resonance overlap stability criterion} \\citep{mar07}. Because of the chaotic nature induced by the nonlinear overlap of several orbital mean motion resonances, the fluid elements inside the chaotic regions significantly change their orbits on a dynamical timescale and do not return to the tidal radius to fuel the BH, leading to the interruption of accretion. We also investigated the evolution of the stellar debris using three-body scattering experiments. Our results obtained both analytically and numerically show that the accretion rate of the debris gas decreases with a power law $\\dot{M} \\sim t^{-5/3}$ until a critical time $T_{\\rm tr}$. At time $t>T_{\\rm tr}$, the accretion pauses until about one SMBHB orbital period $T_{\\rm b}$. $T_{\\rm tr}$ relates to $T_{\\rm b}$ with $T_{\\rm tr} = \\eta T_{\\rm b}$, where $\\eta$ is typically $0.25$ and in the range $0.15 \\la \\eta \\la 0.5$, depending on the initial orbital parameters ($\\theta, \\Omega, \\omega$) of the fluid elements at tidal radius and on the eccentricity of the SMBHB, but being nearly independent of the SMBHB semimajor $a_{\\rm b}$, BH masses, and mass ratio $q$. This suspension of accretion would result in an interruption of the tidal flare, although the residual accretion disk may still radiate weak optical, UV, and X-ray emission during the interruption.  Our numerical results indicate that the accretion of the gaseous stellar debris restarts at a time $T_{\\rm  r}$,  leading to a flicker of flare. The exact recurring time, $T_{\\rm r} = \\xi T_{\\rm b}$, depends on the initial orbital parameters $\\theta$, $\\Omega$, and $\\omega$, but our numerical results suggest that $\\xi$ is of order unity and $1 \\la \\xi \\la 1.5$. The interruption timescale of the tidal flare in a SMBHB system is \\begin{equation} \\Delta{T} = T_{\\rm r} - T_{\\rm tr} \\simeq (\\xi - \\eta) T_{\\rm b} . \\label{int_timescale} \\end{equation} Our numerical simulations suggest that the accreted plasma during the discrete accretion consists of both the tidal debris gas falling back to the tidal radius after one Keplerian time and a fraction of those fluid elements with chaotic orbits falling back to the accreting torus on a timescale longer than its Keplerian time. When $T_{\\rm tr}$ and $\\Delta{T}$ is determined observationally, constraints on $\\eta$ and $T_{\\rm b}$ could be made if we take $\\xi \\simeq 1$.  In our simulations, we assume that the orbital binding energy $E_b$ of the tidally disrupted star is negligible compared to the spread in specific energy $\\Delta E$. For such tidal-disruption events, Equations~(\\ref{eq:tmin}) and (\\ref{tr_th}) imply that the standard power-law decay and the interruption of tidal flares are detectable if \\begin{eqnarray} \\beta < \\beta_{\\rm max} &\\approx& 130 \\sigma_{110}^{-2} M_7^{1/3} q_{-1} (1+q)^{-4/3} \\left({k\\over2}\\right) \\nonumber\\\\ && \\left(\\eta\\over0.25\\right)^{2/3} \\left({R_* \\over \\rsun}\\right)^{-1} \\left({m_{*} \\over \\msun}\\right)^{2/3} . \\label{beta_lim} \\end{eqnarray} If the tidally disrupted star is initially very bound to one of the binary BHs so that $|E_b|$ is comparable to $\\Delta E$ \\citep{che09}, the time $T_{\\rm min}$ would be much shorter and the interruption of tidal flare is detectable even in an ultra-hard SMBHB with $\\beta\\gg\\beta_{\\rm max}$. For a SMBHB with $\\beta \\sim 100$, the semi-major axis of the binary $a_{\\rm b}$ is about $a_{\\rm b} \\simeq 9.3 \\times 10^{2} r_{\\rm g} q_{-1} (1+q)^{-1} \\sigma_{110}^{-2}$. The GW radiation emitted by such SMBHBs could be detected by PTA.  An interrupted tidal flare could be caught if $T_{\\rm tr}$ is shorter than the mission duration of an transient survey $t_{\\rm sv} \\sim 1-10 \\, {\\rm yr}$, which corresponds to SMBHBs with $\\beta \\sim (5 - 100) \\times q_{-1} M_7^{1/6}$. Here we use the $M_\\bullet$--$\\sigma_*$ relation \\citep{tre02}. Because a hard SMBHB with $q=0.1$ spends most of its life time at $\\beta\\sim 5-100$ \\citep{yu02,ses08}, we have a good chance to detect them with upcoming transient surveys. However, if one wants to catch an interrupted tidal flare from SMBHBs emitting strong GW radiation ($\\beta\\sim 1000$), the time resolution of the survey should be $\\la 0.1 \\, {\\rm yr}$. Because a SMBHB spends a small fraction of its lifetime at this stage, high-sensitivity and deep transient surveys are needed to accumulate many more tidal-disruption events."
    },
    "0910/0910.4364_arXiv.txt": {
        "abstract": "The physics items presented in the HE 2.2-2.5  sessions  include a variety of results concerning neutrino oscillations, dark matter and anti-matter, supernova neutrinos and proton decay. The OG 2.5-2.7 sessions concern detector R\\&D and operations in gamma and neutrino astronomy. I report here about a selection of the presented results. These indicate that Neutrino Astronomy is in a mature phase. Challenging detectors are being operated and produce results of interest. Their sensitivity is now getting closer and closer to the region of interest for various predictions. For galactic sources, if expected neutrino fluxes are derived from measured gamma fluxes, it is possible that a few years of operation of cubic-km detectors are needed to produce a statistically significant discovery. The era of the cubic-km detectors is now a reality in the Antarctic ice. The Mediterranean sea and Lake Baikal are hosting smaller scale detectors and proposing future extensions to the cubic-kilometer scale. Long baseline experiments and Super-Kamiokande are measuring with increasing precision the neutrino mixing matrix elements while waiting for dedicated experiments to measure $\\theta_{13}$. Indirect detection of dark matter (DM) with gammas and neutrinos covers a complementary parameter space than LHC and direct detection experiments. Even if these measurements are affected by astrophysical uncertainties, they can provide a hint on the nature of the dark matter particle. Anti-matter limits on deuterons and anti-protons from PAMELA and balloon experiments, such as BESS-Polar, are of high interest and AMS-02 will extend the sensitivity to higher energy and lower fluxes. While past conferences (such as ICRC2003) where dominated by MeV-GeV neutrino results on atmospheric and solar oscillations, this conference was dominated by gamma astronomy results. The field is rich of R\\&D programs that are especially intriguing for what concerns photo-detectors, such is the case of Geiger-APDs. Several works showed that there is room to increase sensitivities of Imaging Atmospheric Cherenkov detectors by not negligible factors with more clever analysis tools and topological triggers. ",
        "introduction": "\\label{sec1} The sessions I have to summarize concern: HE 2.2 Observations on Solar and Atmospheric Neutrinos (13 contributions); HE 2.3 Search for new particles and phenomena (as dark matter,...) (30 contributions); HE 2.4 New Experiments and Instrumentation (25 contributions); OG 2.5 High Energy Neutrino Astrophysics (42 contributions); OG 2.6 Gravitational waves and Experiments (no contributions); OG 2.7 New Experiments and Instrumentation (83 contributions).  Most of the works are about experimental results of detectors in operation and a lot of R\\&D studies especially for Imaging Cherenkov Atmospheric detectors. There were only a few theory contributions. The gravitational wave (GW) community was absent at this conference. It is possible that this is a sign that there is not yet enough exchange between the GW and astroparticle communities. This coming decade may be dominated by the effort of opening a new window on the universe using the yet unobserved GWs. Their detection will test General Relativity and help understanding astrophysical sources. The next decade may as well see the discovery of the other longest range messengers, astrophysical neutrinos. Moreover, ideally the discovery of neutrino and gravitational sources should be accompanied by electro-magnetic observations. The interconnections between GW, neutrino and gamma astronomy are so strong that multi-wavelength approaches are essential to provide the most complete possible picture of astrophysical sources and scientific communities working in these fields should hopefully increase their level of interaction.  It is very difficult for me to find a guiding line between all of these different subjects to review. They share one aspect in my view: they address the question on the mass/energy content of the universe. This paper is structured as follows: in Sec.~\\ref{sec2}  I will address the question: what are we still learning on the problem of neutrino mass through oscillation measurements? In Sec.~\\ref{sec3}  I will discuss upper limits on supernova collapse and proton decay. In Sec.~\\ref{sec4} I will describe what was presented about the mass content in the Universe and its nature, specifically about  indirect searches for DM using gammas and neutrinos and on anti-matter. In Sec.~\\ref{sec5} I will focus on searches for astrophysical sources in neutrino telescopes  focusing on point-like sources and diffuse fluxes. In Sec.~\\ref{sec6} I will describe some of the R\\&D programs and proposed experiments in the neutrino  and gamma astronomy fields. ",
        "conclusions": "\\label{sec7}  Most of the results presented at this conference on neutrino astronomy and DM are upper limits, even if IceCube has showed point-source results for 6 months of data taken with 1/2 of its full configuration (40 strings). These limits are beginning to be in the region of interest between $E^2 \\frac{dN}{dE} \\sim 10^{-12}-10^{-11}$ TeV cm$^{-2}$ s$^{-1}$ for galactic sources and diffuse flux limits are approaching the Waxman \\& Bahcall upper limit on extra-galactic neutrino fluxes.  According to estimates based on gamma observations, it is possible that $\\sim 5$ years of IceCube will be needed to achieve $5\\sigma$ significance for Milagro Pevatrons \\cite{halzen_pevatrons}. Another detector in the opposite hemisphere would cover the entire sky and be sensitive to the Galactic Centre region, while IceCube has sensitivity to Cygnus that is observed by Milagro up to tens of TeV. Studies and R\\&D for detectors of the order of 10-100 times IceCube are a promising strategy for the future, but no technique is yet available at a reasonable cost for the same energy threshold as current NTs. In fact, radio techniques that are very promising to achieve 100 km$^2$, have thresholds of the order of $10^{18}$ eV suitable for cosmogenic neutrinos but considerably higher than IceCube.  Gamma astronomy future programs aim at pushing the energy threshold as low as possible to fill the gap from Fermi-GLAST and possibly observe close-by gamma-ray bursts with ground based arrays in the TeV range. In the future, arrays of IACT telescopes can be suitable for this scope joining the European and American communities. R\\&D for camera detectors should be strongly encouraged. In the meanwhile large FoV arrays with 100\\% duty cycles could monitor the TeV sky complementing at higher energies Fermi-GLAST."
    },
    "0910/0910.0771_arXiv.txt": {
        "abstract": "Dark matter annihilation is one of the leading explanations for the recently observed $e^\\pm$ excesses in cosmic rays by PAMELA, ATIC, FERMI-LAT and HESS. Any dark matter annihilation model proposed to explain these data must also explain the fact that PAMELA data show excesses only in $e^\\pm$ spectrum but not in anti-proton. It is interesting to ask whether the annihilation mode into anti-proton is completely disallowed or only suppressed at low energies. Most models proposed have negligible anti-protons in all energy ranges. We show that the leptocentric $U(1)_{B-3L_i}$ dark matter model can explain the $e^\\pm$ excesses with suppressed anti-proton mode at low energies, but at higher energies there are sizable anti-proton excesses. Near future data from PAMELA and AMS can provide crucial test for this type of models. Cosmic $\\gamma$ ray data can further rule out some of the models. We also show that this model has interesting cosmic neutrino signatures. ",
        "introduction": "Recently several experiments have reported $e^\\pm$ excesses in cosmic ray (CR) energy spectrum. Last year the PAMELA collaboration reported $e^+$ excess in the CR energy spectrum from 10 to 100 GeV, but observed no anti-proton excess~\\cite{pamela-e,pamela-p} compared with predictions from CR physics. These results are compatible with the previous HEAT and AMS01 experiments (e.g., Ref. \\cite{heat,ams}) but with higher precision. Shortly after the ATIC and PPB-BETS balloon experiments have reported excesses in the $e^+ + e^-$ spectrum between 300 and 800 GeV~\\cite{atic,ppb-bets}. The ATIC data show a sharp falling in the energy spectrum around 600 GeV. Newly published result from FERMI-LAT collaboration also shows excesses in the $e^+ + e^-$ energy spectrum above the background~\\cite{fermi}. However, the spectrum is softer than that from ATIC. In addition, the HESS collaboration has inferred a flat but statistically limited $e^++e^-$ spectrum between 340~GeV and 1~TeV~\\cite{hess1} which falls steeply above 1~TeV~\\cite{hess2}. These observational data have generated a lot of excitement as the excesses may be explained by dark matter (DM) annihilation or decay with appropriate properties \\cite{darkmatter,df,lepto,bi-he,review}.  If DM annihilation is responsible for the observed data, a lot of information can be extracted about DM properties\\cite{darkmatter,df,lepto,bi-he,review}. The mass of the annihilating DM serves as the cut off scale of the $e^\\pm$ spectrum, the lepton spectra must have a cut off energy at the DM mass $m_D$. The FERMI-LAT and HESS data would require that the DM mass to be around  a few TeV. The DM belongs to the weakly interacting massive particle (WIMP) category. The fact that PAMELA did not see anti-proton excess indicates that the DM is hadrophobic or leptophilic. At least hadronic annihilation modes are suppressed. Also since the same DM annihilation rate producing the $e^\\pm$ excesses is related to the annihilation rate producing the cosmological relic DM density, 23\\% of the energy budget of our universe, the annihilation rate is constrained. The latter requires that the thermally averaged annihilation rate $\\langle \\sigma v \\rangle$ to be $3\\times 10^{-26}$ cm$^3$s$^{-1}$. If this annihilation rate is used to calculate the $e^\\pm$ spectrum, it is too small by a factor of 100 to 1000. There is the need to boost up the spectrum with a boost factor $B$ of 100 to 1000. Any DM model proposed to explain the data should be able to produce the large boost factor too.   To suppress the cosmic anti-protons leptophilic DM models are proposed by several authors \\cite{lepto,bi-he}. In this type of models the DM only interacts with leptons in the standard model (SM). However, these DM models are hard to be tested since DM does not interact with hadrons at all. Another widely considered $U(1)$ extension of the SM is the gauged B-L. However, if DM interacts with the SM with the gauged $U(1)_{B-L}$ interaction too much anti-protons will be produced and therefore in conflict with PAMELA data. In this work, we propose a particle physics DM model, the gauged~\\cite{b-3t} $U(1)_{B-3L_i}$ DM model, to explain the $e^\\pm$ excess data. Here B is the baryon number and $L_i$ is one of the $e$, $\\mu$ and $\\tau$ lepton numbers.  In this model, the DM is a Dirac fermion with nontrivial $U(1)_{B -3 L_i}$ charge~\\cite{b-3t}. The gauge boson of this $U(1)$ group, $Z^\\prime$, mediates DM annihilation into SM particles. This model is not a pure leptophlic DM model, like the $U(1)_{L_i-L_j}$ model studied in Ref. \\cite{bi-he}. But has larger couplings to leptons than quarks. We refer it as the leptocentric dark model. The $Z^\\prime$ has suppressed couplings to quarks leading to suppressed anti-proton up to the PAMELA energy reach, but may lead to excesses at higher energies. With the $Z^\\prime$ mass close to two times of the DM mass, a large boost factor can be produced through the Breit-Wigner enhancement mechanism~\\cite{df,breit-wigner}. The shape of the $e^\\pm$ spectrum helps to further constrain the choice of the lepton flavor $L_i$ of the $U(1)_{B-3 L_i}$.   All DM annihilation models proposed to explain $e^\\pm$ contribute at certain level to $\\gamma$-rays by final state radiation. One must check if the model is compatible with available data. We find that for the $L_i = \\tau$ model, the hadronic $\\tau$ decays lead to too much $\\gamma$-rays in conflict with data and only the $L_i = L_{e,\\mu}$ models can pass the $\\gamma$ ray data constraint if one takes the $e^\\pm$ excess spectrum to fit the ATIC or FERMI-LAT data. It is interesting to note that the $Z^\\prime$ has the same couplings for charged lepton $l_{L_i}$ and neutrino $\\nu_{L_i}$. Therefore after fitting data on $e^\\pm$ excesses in CR, the cosmic neutrinos by DM annihilation from the Galactic center (GC) are predicted. The neutrino signals can be tested by experiments such as IceCube.  In the next section we will first describe how to treat the propagation of charged cosmic rays in the Milky Way. Then we present the leptocentric $U(1)_{B-3L_i}$ DM model in Sec. III. The cosmic $e^\\pm$, $\\bar{p}$, $\\gamma$ and neutrino spectra predicted are given in Sec. IV, V, VI and VII respectively. Finally we give our conclusions in Sec. VIII. ",
        "conclusions": "In this work we studied cosmic  $e^\\pm$, $\\bar p$, $\\gamma$ and neutrino rays in the leptocentric $U(1)_{B-3 L_i}$ dark matter models. In these models, DM annihilation into SM particles is mediated by the  $Z^\\prime$ gauge boson of $U(1)_{B-3 L_i}$. The couplings of $Z^\\prime$ to leptons are larger than those to quarks leading to larger annihilation of DM to leptons than hadrons. This naturally explains the $e^\\pm$ excesses from PAMELA, ATIC or FERMI observational data, and at the same time these models can suppress the anti-proton flux efficiently. We find that the $U(1)_{B-3 L_e}$ model can fit PAMELA and ATIC, while $U(1)_{B-3 \\mu}$ and $U(1)_{B-3 \\tau}$ can fit PAMELA and FERMI, respectively. Future data can be used to distinguish different models. We showed that by slightly adjusting the propagation parameters these models can predict anti-proton flux consistent with the PAMELA data. But these models are different than pure leptophilic models which predict negligible anti-proton cosmic ray in all energy ranges, the leptcentric $U(1)_{B-3L_i}$ models predict anti-proton flux beyond the background at higher energies which can be tested by the near future data from PAMELA or AMS02.  The $U(1)_{B-3L_i}$ models have predictions for cosmic $\\gamma$-rays from DM annihilation. We investigated the $\\gamma$-ray emission from the GC in these models. We found that the  $U(1)_{B-3 \\tau}$ model is excluded by the HESS observation of the GC region since this model predicts too much $\\gamma$-rays. Adopting the Einasto DM density profile the other two models predict $\\gamma$-rays consistent with HESS data.  We also investigated the detectability of neutrinos from the GC by the $U(1)_{B-3 e}$ and $U(1)_{B-3 \\mu}$ models at DeepCore/IceCube. In these models the fraction of the DM annihilation into neutrinos are large and produce line spectrum neutrinos detectable by DeepCore/IceCube. We found that the DeepCore/IceCube can have a $8.5(2.6)\\sigma$ signal of the line neutrino spectrum for 4 years data taking. This will be a distinctive signal of DM annihilation in our leptocentric $U(1)_{B-3e}$ and $U(1)_{B-3\\mu}$ models.  Finally we would like to have some comments on the detectability of leptocentric DM model at the LHC and direct DM detection experiments. Since the leptocentric DM models discussed here are very different than the pure leptophilic DM models, where the $Z^\\prime$ only has interactions with leptons at the tree level, the $Z^\\prime$ in our cases has interaction with quarks and may lead to detectable effects at the LHC and also direct DM search on the Earth. However because we rely on Breit-Wigner resonant enhancement effect to produce the large boos factor, the $Z^\\prime$  mass is constrained to be 2 times of the DM mass and is constrained to be in the TeV range and the couplings $a g{^\\prime}^2$ to be in the range of $1\\times10^{-5}$ to $5\\times10^{-5}$, the cross section for producing $Z^\\prime$ and virtual effects due to $Z'$ for the LHC and also for direct detection are small, it is not possible to have observable effects at the LHC and near future direct DM detection experiments. To have a lower $Z^\\prime$ mass and therefore to have observable effects at the LHC and direct DM detection experiments, a different mechanism for the boost factor has to be in effective to relax the relation that the $Z^\\prime$ mass is about 2 times of the DM mass.  To this end we note that if there is a long range interaction between DM by exchanging a light new particle, it is possible to have the Sommerfeld effect in operation, e.g. in Ref. \\cite{review,hisano}. This offers another interesting possibility to explain the cosmic data. We will give details for this possibility elsewhere."
    },
    "0910/0910.5362_arXiv.txt": {
        "abstract": "{} { We investigate the stellar pancake mechanism during which a solar-type star is tidally flattened within its orbital plane passing close to a $10^{6} \\, M_{\\odot}$ black hole. We simulate the relativistic orthogonal compression process and follow the associated shock waves formation. } { We consider a one-dimensional hydrodynamical stellar model moving in the relativistic gravitational field of a non-rotating black hole. The model is numerically solved using a Godunov-type shock-capturing source-splitting method in order to correctly reproduce the shock waves profiles. } { Simulations confirm that the space-time curvature can induce several successive orthogonal compressions of the star which give rise to several strong shock waves. The shock waves finally escape from the star and repeatedly heat up the stellar surface to high energy values. Such a shock-heating could interestingly provide a direct observational signature of strongly disruptive star - black hole encounters through the emission of hard $X$ or soft $\\gamma$ -ray bursts. Timescales and energies of such a process are consistent with some observed events such as GRB 970815. } {} ",
        "introduction": "It has long been underlined that stars orbiting closely enough to massive black holes could be tidally flattened into a transient \\textit{pancake}-shape configuration until full disruption \\citep{car1982,car1983}. Strongest compressions have been predicted to trigger  thermonuclear explosion in the core of small main-sequence stars grazing black holes from thousands up to millions solar masses  \\citep{pic1985,lum1989a}. The presence of specific proton-enriched chemical elements near galactic centres may be compatible with the stellar pancakes nucleosynthesis \\citep{lum1990}.  More recently, high-resolution hydrodynamical simulations have shown that the tidal compression could independently give rise to strong shock waves within the stellar matter \\citep{kob2004,bra2008,gui2009}. Propagating outwards, shock waves may then be able to quickly heat up the stellar surface and lead to the emission of a new type of hard $X$ or soft $\\gamma$ -ray bursts.  Aftermath of tidal disruptions of stars by massive black holes have already been detected from nearby otherwise quiescent galactic cores through giant-amplitude persistent $X$-$UV$ flares \\citep[e.g.][]{bad1996,gru1999,kom1999,gre2000,kom2002,hal2004,kom2004,gez2006,esq2007,gez2008,cap2009}. These flares have arised from the long term evolution when part of the liberated stellar gas was being accreted by the central black holes. However, the very short-timescale high-energy flares due to the initial compression of the star could interestingly allow to catch disruptions from their beginning.  In our previous article \\citep[][hereafter \\citetalias{bra2008}]{bra2008}, we have simulated the tidal compression orthogonal to the orbital plane using a one-dimensional Godunov-type shock-capturing hydrodynamical model, and assuming that the star evolved within a Newtonian external gravitational field.  In this article, we naturally extend the model to the general relativistic case of a non-rotating Schwarzschild black hole and present some complementary results.  In the following, relativistic equations are expressed in geometrized units $c=G=1$, greek and latin indices run respectively from 0 to 3 and 1 to 3. ",
        "conclusions": "We have presented first explorative results on the star's tidal compressions by a massive Schwarzschild black hole which take account of the intrinsic shock waves formation orthogonal to the orbital plane. Whereas the stellar core would be heated up by the external compressive tidal field (Fig.~\\ref{fig_i_01}), the hydrodynamical simulations also suggest that the stellar surface would be heated up during the shock waves propagation. \\begin{figure} \\begin{center} \\begin{tabular}{c} \\includegraphics[width=8.5cm]{pf_rhocmax_1.eps} \\\\ \\includegraphics[width=8.5cm]{pf_tempcmax_1.eps} \\\\ \\includegraphics[width=8.5cm]{pf_timecmax_1.eps} \\end{tabular} \\caption{ Evolutions of the maximum central density (top), maximum central temperature (middle), and duration during which the highest values are maintained (bottom) as functions of the penetration factor $\\beta$ in the case $\\gamma=5/3$. The rightside up (resp. upside down) triangles indicate the first (resp. second) compression. The points show for comparison the single compression calculated for a Newtonian gravitational field, and the dashed lines correspond to the power laws predicted by the affine model (\\citealt{car1983}; see also \\citetalias[][Eqs.~(39)-(41) and Fig. 16]{bra2008}). In bottom figure, the characteristic internal timescale of the star $\\tau_{\\star}=(1/G\\rho_{\\star})^{1/2}$. } \\label{fig_i_01} \\end{center} \\end{figure}  Besides the more restrictive scenario of the core thermonuclear explosion, it is thus believed that the close passage of a star to a massive black hole may be directly revealed as highly energetic short-timescale bursts originating from the shocked stellar surface. It is well known that some observed gamma-ray light curves present two or more peaks with complex structures \\citep[e.g.][]{fis1995}. Could some of them be interpreted as tidally-induced pancake stars? Typical timescales shown in Fig.~\\ref{fig_g_10} for a star - black hole encounter along a relativistic self-crossing orbit are quite consistent with those of a short duration gamma-ray burst with two  peaks of energy separated by a hundred of seconds, comparable to what has been observed e.g. in GRB 970815 \\citep{smi2002}. The long duration GRB 060614 has also been tentatively interpreted by other authors as a strongly disruptive star - black hole encounter \\citep{lu2008}.  Tidally-induced gamma-ray bursts are estimated to occur once every $10^{3}$--$10^{5}$ years per galaxy, depending on the nuclear stellar density profile and the mass of the central black hole \\citep{wan2004}. Such a rate can even be increased by a factor $10^{4}$ in the case of a massive binary black hole \\citep{che2009}. Since most of the galaxies --~including our own Milky Way~-- harbour a central massive black hole, and since the full observable universe is transparent in $X$ or $\\gamma$ -ray domains, several events of this kind would be then detectable each year.  To well characterize this signature and make valuable comparisons with observational data of high-energy bursts, simulations have to consider the full stellar pancake's structure and the radiative transfer of energy through the layers of the star. From this perspective, three-dimensional high-resolution shock-capturing simulations have been recently performed in the context of a Newtonian gravitational field, and have shown that shock waves actually form in several other directions within the star with the consequence of isotropizing the surface temperature \\citep{gui2009}."
    },
    "0910/0910.2068_arXiv.txt": {
        "abstract": "We report the results of CCD $V$, $r$ and $I$ time-series photometry of the globular cluster NGC 5053. New times of maximum light are given for the  eight known RR Lyrae stars in the field of our images and their periods are revised. Their $V$ light curves were Fourier decomposed to estimate their physical parameters. A discussion on the accuracy of the Fourier-based iron abundances, temperatures, masses and radii is given. New periods are found for the 5 known SX Phe stars and a critical discussion of their secular period changes is offered. The mean iron abundance for the RR Lyrae stars is found to be [Fe/H] $\\sim -1.97 \\pm 0.16$ and lower values are not supported by the present analysis. The absolute magnitude calibrations of the RR Lyrae stars yield an average true distance modulus of $16.12 \\pm 0.04$ or a distance of $16.7 \\pm 0.3$ kpc. Comparison of the observational CMD with theoretical isochrones indicates an age of  $12.5 \\pm 2.0$ Gyrs for the cluster. A careful identification of all reported Blue Stragglers (BS) and their $V,I$ magnitudes leads to the conclusion that BS12, BS22, BS23 and BS24 are not BS.  On the other hand, three new BS are reported.  Variability was found in seven BS, very likely of the SX Phe type in five of them, and in one red giant star. The new SX Phe stars follow established $PL$ relationships and indicate a distance in agreement with the distance from the RR Lyrae stars. ",
        "introduction": "Recent numerical methods in the analysis of stellar CCD images have proven to be very powerful in isolating faint stars and performing accurate photometry even in very crowded fields, such as in the central regions of globular clusters (Bramich 2008, Bramich et al. 2005; Alard 2000; Alard \\& Lupton 1998). In a series of papers we have taken advantage of these mathematical techniques to study the known variable stars and to search for new ones in several globular clusters of assorted metallicities (Arellano Ferro et al. 2008a; 2008b; 2004, L\\'azaro et al. 2006). In particular we used the Fourier decomposition technique to estimate fundamental physical parameters of the RR Lyrae stars and we searched for eclipsing binaries and variability among the blue stragglers. The calculation of the iron content [Fe/H] and absolute magnitude $M_V$ lead to a linear $M_V$-[Fe/H] relationship (Arellano Ferro et al. 2008b) that is in agreement with independent estimations from different empirical techniques (e.g. Cacciari \\& Clementini 2003; Chaboyer 1999).  The globular cluster NGC 5053 (R.A.(J2000)$=13^h16^m27^s$.3, DEC(J2000)$=+17^{\\circ} 41'52''$, $l=335.7$, $b=+78.9$) is located in the intermediate Galactic halo ($Z=$ 16.1 kpc, R$_G$=16.9 kpc) and it is one of the most metal deficient clusters; [Fe/H] = $-2.1$ (Rutledge et al. 1997). It is a rather disperse and well resolved system which allows high quality photometry even in the central regions. Such characteristics place this stellar system close to the grey region between globular clusters and open clusters. Nevertheless it was classified as a globular cluster mainly due to its high galactic latitude and the presence of faint stars and some variables (Cuffey 1943, Baade 1928). The cluster has been the subject of detailed photometric studies and the most recent ones, already in the CCD era, conducted detailed studies of known RR Lyrae stars (Nemec, Mateo \\& Schombert 1995) and the identification of 28 blue stragglers (Nemec \\& Cohen 1989; Sarajedini \\& Milone 1995) among which 5 are known to be SX Phe variables (Nemec et al. 1995).  In the present paper we perform standard $V$, $I$ and instrumental $r$ CCD photometry for nearly 6500 stars in the field of NGC 5053 and perform the Fourier light curve decomposition of the known RR Lyrae stars. We also search for variability among Blue Straggler (BS) stars and revisit the known SX Phe stars previously studied by Nemec et al. (1995), in an attempt to refine their periodicities in order to discuss their period changes and to use them as independent indicators of the distance to the cluster. Finally we report five new SX Phe stars, three new BS stars, and one new red giant variable (RGV).  The paper is organised as follows: in $\\S$ 2 we describe the observations, data reductions and transformations to the standard photometric system. In $\\S$ 3 a period analysis of the RR Lyrae stars is performed. In $\\S$ 4 the $V$ light curves of the RR Lyrae stars are Fourier decomposed and their physical parameters are estimated. In $\\S$ 5 the periods and secular period changes of the SX Phe stars are discussed as well as the BS nature of previously reported BS stars. In $\\S$ 6 a search for new variables is carried out and the new variables are described. In $\\S$ 7 the distance and age of NGC~5053 are discussed and in $\\S$ 8 we summarise our conclusions. ",
        "conclusions": "The technique of difference imaging has proven to be a powerful tool in providing accurate photometry down to about $V=20$ mag and in finding new variables. In this paper we report the discovery of five SX Phe stars among the BS stars, two of which are new BS stars, and one new red semiregular variable.  Physical parameters of astrophysical relevance such as: $log~(L/L_{\\odot})$, $log~T_{\\rm eff}$, $M_V$, $R/R_{\\odot}$, $M/M_{\\odot}$ and [Fe/H], have been derived for the RR Lyrae stars in NGC~5053 using the Fourier decomposition of their $V$ light curves. Special attention has been paid to the calibration of these parameters onto widely accepted scales. It has been found that $T_{\\rm eff}$ for the RRc stars from the calibration of Simon \\& Clement (1993) cannot be reconciled with the theoretical predictions for the boundaries of the instability strip. However, the Fourier temperatures for RRab stars from the calibration of Jurcsik (1998) agree well with the colour temperatures, and they imply radii and masses comparable to those derived from pulsational models.  The average iron abundance [Fe/H] = $-1.97 \\pm 0.16$ is found from the RRab and RRc stars. This value may appear too large for a cluster often considered the most metal deficient in the Galactic halo. However, this value is in good agreement, within the uncertainties, with the most recent accurate estimates from spectroscopic data.  A distance of $16.7 \\pm 0.3$ kpc is found from the Fourier approach to the determination of $M_V$ for the RR Lyrae stars. This value is in agreement, within uncertainties, with predictions from the $M_V$-[Fe/H] relation for values of [Fe/H] $\\sim -2.0$.  We report five new SX Phe stars and show that they follow the period luminosity relationship of SX Phe stars in NGC~5466, a cluster very similar in age and metallicity to NGC~5053. If the $PL$ calibration is applied to the SX Phe in NGC~5053, then the estimated distance is found in good agreement with the distance from the RR Lyrae stars.  Our $(V-I)$ CMD and the isochrones of VandenBerg, Bergbusch \\& Dowler (2006) indicate an age of $12.5 \\pm 2.0$ Gyrs for the cluster, which is in good agreement with independent differential age estimations."
    },
    "0910/0910.5224_arXiv.txt": {
        "abstract": "\\bao~(BAO) are frozen relics left over from the pre-decoupling universe. They are the standard rulers of choice for 21st century cosmology, providing distance estimates that are, for the first time, firmly rooted in well-understood, linear physics. This review synthesises current understanding regarding all aspects of BAO cosmology, from the theoretical and statistical to the observational, and includes a map of the future landscape of BAO surveys, both spectroscopic and photometric\\footnote{This review is an extended version of a chapter in the book {\\em Dark Energy} Ed.~Pilar Ruiz-Lapuente, Cambridge University Press (2010, ISBN-13: 9780521518888)}. ",
        "introduction": "Whilst often phrased in terms of the quest to uncover the nature of dark energy, a more general rubric for cosmology in the early part of the 21st century might also be the ``the distance revolution\". With new knowledge of the extra-galactic distance ladder we are, for the first time, beginning to accurately probe the cosmic expansion history beyond the local universe. While standard candles -- most notably Type Ia supernovae (SNIa) -- kicked off the revolution, it is clear that Statistical Standard Rulers, and the \\bao~(BAO) in particular, will play an  increasingly important role.  In this review we cover the theoretical, observational and statistical aspects of the BAO as standard rulers and examine the impact BAO will have on our understanding of dark energy, and the distance and expansion ladder. Fisher matrix forecasts for BAO surveys can be easily computed using the publically released, GUI-based, \\name~code, which is publically available\\footnote{http://www.cosmology.org.za}. \\subsection{A Brief History of Standard Rulers and the BAO} Let us start by putting the BAO in context. The idea of a standard ruler is one familiar from everyday life. We judge the distance of an object of known length (such as a person) by its angular size. The further away it is, the smaller it appears. The same idea applies in cosmology, with one major complication: space can be curved. This is similar to trying to judge the distance of our known object through a smooth lens of unknown curvature. Now when it appears small, we are no longer sure it is because it is far away. It may be near and simply appear small because the lens is distorting the image. This degeneracy between the curvature of space and radial distance has not been the major practical complication in cosmology over the past century, however. That honour goes to a fact that has plagued us since the beginning of cosmology: we don't know how big extragalactic objects are in general, in the same way that we don't know how bright they intrinsically are. This problem was at the heart of the great debate between Shapley and Curtis over the nature of galaxies. Shapley argued that they were small and inside our own galaxy while Curtis maintained that they were extragalactic and hence much larger.  To be useful for cosmology, we need a {\\em standard ruler}: an object of a known size at a single redshift, $z$, or a population of objects at different redshifts whose size changes in a well-known way (or is actually constant) with redshift. Ideally the standard ruler falls into both classes, which, as we will argue below, is the case for the BAO, to good approximation. BAO are however a new addition to the family of putative standard rulers. A few that have been considered in the past include ultra-compact radio sources \\cite{kellermann93,gurvits94} which indeed lead in 1996, prior to the SNIa results, to claims that the density of dark matter was low, $\\Omega_m < 0.3,$ with a non-zero cosmological constant of indeterminate sign \\cite{jackson97}. Another radio standard ruler candidate is provided by double-lobed radio sources \\cite{buchalter_radio}. These Fanaroff-Riley Type IIb radio galaxies were suggested as cosmological probes as early as 1994 \\cite{daly_friib}, and subsequent analyses have given results consistent with those from SNIa \\cite{daly, latest_daly}. An alternative approach uses galaxy clusters. Allen {\\em et al.} relate the X-ray flux to the cluster gas mass, and in turn, its size, providing another standard ruler under the assumption that the gas fraction is constant in time. This too leads to results consistent with those from SNIa \\cite{allen_schmidt, latest_allen}.  Beyond this we move into the realm of Statistical Standard Rulers (SSR), of which BAO are the archetype. SSR exploit the idea that the clustering of galaxies may have a preferred scale in it which, when observed at different redshifts, can be used to constrain the angular diameter distance. The idea of using a preferred clustering scale to learn about the expansion history of the cosmos has a fairly long history in cosmology, dating back at least to 1987 and perhaps earlier. In their conclusions Shanks {\\em et al.} \\cite{shanks87} forsee: \\begin{quote} {\\em There is one further important reason for searching for weak features in $\\xi_{qq}(r),$ the quasar-quasar correlation function, at large separations. If a particular feature were found to appear in both the galaxy correlation function at low redshift and the QSO correlation function at high redshift, then a promising new cosmological test for $q_0$ might be possible.} \\end{quote} A series of later analyses further built up the idea of SSR in cosmology using variously as motivation the turn-over in the power spectrum due to the transition from radiation to matter-domination, the mysterious $128 h^{-1} \\mathrm{Mpc}$ feature detected in early pencil-beam surveys \\cite{broadhurst_ellis} and the realisation that inflation could inject a preferred scale into the primordial power spectrum. These early studies often found tentative evidence for a low-density universe and/or non-zero cosmological constant e.g. \\cite{deng,broad_jaffe_ssr}. Since the preferred scale could not be accurately predicted {\\em a priori} these studies only provided the relative size of the SSR at different redshifts. Nevertheless, it was realised that this could provide interesting constraints on the expansion history of the universe \\cite{roukema9911413,roukema2_0106135}.  BAO entered the fray initially as a putative explanation for the apparent excess clustering around $100 h^{-1} \\mathrm{Mpc}$  but were found to be too weak to be the origin for the apparent excess \\cite{eisenstein_hu_1998, meiskin}. The idea of using BAO themselves to learn about cosmological parameters seems to date first from Eisenstein {\\em et al.} (1998) \\cite{eis98_complimentarity} who wrote: \\begin{quote} {\\em Detection of acoustic oscillations in the matter power spectrum would be a triumph for cosmology, as it would confirm the standard thermal history and the gravitational instability paradigm. Moreover, because the matter power spectrum displays these oscillations in a different manner than does the CMB, we would gain new leverage on cosmological parameters.} \\end{quote} The first photometric proposal for using the BAO as standard rulers for learning about cosmology appears to date from 2001 \\cite{cooray_halo}. The real foundations, however, of the modern ideas on BAO, their detection and use, were laid by Eisenstein (2003) \\cite{eisenstein2003}, Blake and Glazebrook  (2003) \\cite{blake_glaze_simulation} and a slew of later papers \\cite{hu_haiman,amendola_de, blake_bao, glaze_blake, white_bao, wang_bao, huff, angulo, komatsu_angular_bao, parkinson06, dunemartin} which developed hand-in-hand with the analysis of real data.  Tantalising hints for the existence of the BAO were already visible in the Abell cluster catalogue \\cite{miller}, but definitive detections had to wait for the increased survey volume and number density of galaxies achieved in the SDSS and 2dF redshift surveys, which immediately yielded strong constraints on both curvature and dark energy at $z < 0.5$ \\cite{eisenstein_05,cole_2df,tegmark_lrg, percival_lrg, hui_bao}. Figures (\\ref{eisenstein:fig}) and (\\ref{sdss:fig}) show the original evidence for the acoustic signature in the correlation function and power spectrum.  Extracting the BAO scale from the matter power spectrum remains a thriving area of research in contemporary cosmology, as we discuss later in Section \\ref{experiments} on current and future BAO surveys. \\begin{figure}[htbp!] \\begin{center} \\includegraphics[width = 3.8in, height = 3.1in]{xi_jack.eps} \\caption{The Baryon Acoustic Peak (BAP) in the correlation function -- the BAP is visible in the clustering of the SDSS LRG galaxy sample, and is sensitive to the matter density (shown are models with $\\Omega_mh^2=0.12$ {\\bf (top)}, 0.13 {\\bf (second)} and 0.14 {\\bf (third)}, all with $\\Omega_bh^2=0.024$). The bottom line without a BAP is the correlation function in the pure CDM model, with $\\Omega_b=0$. From Eisenstein {\\em et al.,}  2005 \\cite{eisenstein_05}.\\label{eisenstein:fig}} \\end{center} \\end{figure}  \\begin{figure}[htbp!] \\begin{center} \\includegraphics[width = 3.8in, height = 3.1in]{all2power.ps} \\caption{Baryon Acoustic Oscillations (BAO) in the SDSS power spectra -- the BAP of the previous figure now becomes a series of oscillations in the matter power spectrum of the SDSS sample. The power spectrum is computed for both the main SDSS sample {\\bf (bottom curve)} and the LRG sample {\\bf (top curve)}, illustrating how LRGs are significantly more biased than average galaxies. The solid lines show the $\\Lambda$CDM fits to the WMAP3 data \\cite{wmap3}, while the dashed lines include nonlinear corrections. Figure from Tegmark {\\em et al.,} 2006 \\cite{tegmark_lrg}. \\label{sdss:fig}} \\end{center} \\end{figure} \\subsection{Cosmological Observables} We now discuss the relevant cosmological observables that are derived from standard rulers in general, and the BAO in particular. The \\bao~in the radial and tangential directions provide measurements of the Hubble parameter and angular diameter distance respectively. The Hubble parameter, $H \\equiv \\dot{a}/a$ -- where $a$ is the scale factor of the universe -- can be written in dimensionless form using the Friedmann equation as \\be E(z) \\equiv \\frac{H(z)}{H_0} = \\sqrt{ \\Omega_m(1+z)^3 + \\Omega_{\\mathrm{DE}}f(z) + \\Omega_k(1+z)^2 + \\Omega_{rad}(1+z)^4}\\,, \\label{hubble_eq} \\ee where $f(z)$ is the dimensionless dark energy density and $\\Omega_k = -\\frac{k}{H_0^2a^2} = 1- (\\Omega_m + \\Omega_{DE} + \\Omega_{rad})$ is the density parameter of curvature with $\\Omega_k = 0$ corresponding to a flat cosmos. $\\Omega_{m},\\Omega_{rad}$ are the matter and radiation densities with corresponding equations of state $w_i \\equiv p_i/\\rho_i = 0,\\frac{1}{3}$ for $i=m,rad$ respectively.  If one treats the dark energy as a barotropic fluid with an equation of state with arbitrary redshift dependence, $w(z)$, the continuity equation can be directly integrated to give the evolution of the dimensionless dark energy density, $f(z)= \\rho_{DE}/\\rho_{DE}(z=0)$, via \\be f(z) = \\exp\\left[3 \\int_0^z  \\frac{1+w(z')}{1+z'}dz'\\right]\\,. \\label{fz} \\ee When we quote constraints on dark energy it will typically be in terms of the CPL parameterisation \\cite{cp,linder_w} \\be w(z) = w_0 + w_a\\frac{z}{1+z}\\,, \\label{wcpl} \\ee which has \\be f(z) = (1+z)^{3(1+w_0+w_a)} \\exp \\left\\{-3w_a \\frac{z}{1+z}\\right\\}\\,.\\label{feq} \\ee  Much of the quest of modern cosmology is to constrain the allowed range of $w(z)$ (or $f(z)$) and hence use this to learn about physics beyond the standard model of particle physics and General Relativity. Apart from direct measurements of the Hubble rate, one of the ways to constrain $w(z)$ using cosmology is through distance measurements. Core to defining distances in the FLRW universe is the dimensionless, radial, comoving distance: \\be \\chi(z) \\equiv\\int_0^z \\frac{dz'}{E(z')}\\,. \\label{chiz}\\ee One then builds the standard cosmological distances using $\\chi(z)$ to give the {angular diameter distance}, $d_A(z)$ via \\be d_A(z) = \\frac{c}{H_0 (1+z) \\sqrt{-\\Omega_k}}\\sin\\left(\\sqrt{-\\Omega_k}\\chi(z)\\right) \\label{da_eq} \\ee and the {luminosity distance}, $d_L(z)$, given via the distance duality as \\be  d_L(z) = (1+z)^2 d_A(z)  \\label{dist_duality} \\ee  \\begin{figure}[htbp!] \\begin{center} \\includegraphics[width = 2in, height = 1.6in]{schematic_curv3.eps} \\caption{Curvature and $\\chi(z)$ define cosmological distances -- In a flat Universe, the cosmological distances are determined by $\\chi(z) \\propto \\int_0^z dz'/E(z').$ In a general FLRW model, however, spatial curvature bends the light rays away from straight lines and hence alters distances, meaning that one needs to know both $\\Omega_k$ and $\\chi(z)$. As a result distance measurements always show a degeneracy between curvature ($\\Omega_k$) and dynamics ($H(z)$). \\label{schematic_dist}} \\end{center} \\end{figure} The expression (Eq.~(\\ref{da_eq})) for $d_A(z)$ holds for all values of the curvature, $\\Omega_k$, since for $\\Omega_k < 0$ the complex argument in Eq.~(\\ref{da_eq}) converts the $\\sin$ function to the $\\sinh$ function. Hence the two key quantities that determine distances in cosmology are the dimensionless distance $\\chi(z)$ and $\\Omega_k$, shown schematically in Figure~(\\ref{schematic_dist}). The link (Eq.~(\\ref{dist_duality})) between $d_A(z)$ and $d_L(z)$ holds in any metric theory of gravity as long as photon number is conserved. This distance duality can be tested and used to look for exotic physics \\cite{bassett_distance2, bassett_distance, uzan_dd, more_dd, verde_dd}.  Distances have a significant disadvantage over pure Hubble measurements: they require an integral over $f(z)$ which is itself an integral over $w(z)$. Hence, any interesting features in $w(z)$ tend to be washed out in distance measurements. There is also another problem: if we look at Eq.~(\\ref{da_eq}) for $d_A(z)$ we notice that if we make no assumptions about $f(z)$, then even perfect distance measurements cannot break the degeneracy between $f(z)$ and $\\Omega_k$ \\cite{weinberg1972}. This is not a fundamental problem if one assumes that $w(z)$ has finite degrees of freedom, e.g. in Eq.~(\\ref{wcpl}), but one must remember that the degeneracy is being broken artificially by hand through ones choice of parameterisation and not by the data. As an example, consider Figure~(\\ref{hda_curv}), which shows two possible dark energy survey configurations; one with $1\\%$ measurements of the angular diameter distance at redshifts of $z=1,3,$ and another with measurements of both the Hubble parameter and angular diameter distance, but with double the errors on each observable. Even given the increase in the error on the observables, the dark energy constraints are significantly improved when including data from these complimentary probes when marginalising over curvature. \\begin{figure}[htbp!] \\begin{center} \\includegraphics[width = 3.8in, height = 3.1in]{h_da_curv.eps} \\caption{Breaking the curvature-dark energy degeneracy -- Error ellipses for the CPL parameters $w_0, w_a$, in two survey scenarios after marginalising over curvature: one consisting of two measurements at $z = 1$ and $z=3$ of the angular diameter distance, with $1\\%$ errors on $d_A(z)$ (blue dashed curve), and another of measurements of $H(z)$ and $d_A(z)$ (solid orange curve) where the errors have been doubled for both observables. In both surveys we assume a prior on curvature of 30, and $\\mathrm{Prior}(\\Omega_m) = \\mathrm{Prior}(H_0) = 1000.$ Even given a weak prior on curvature, combining measurements from multiple probes helps break the curvature-dark energy degeneracy. Figure produced using \\name.\\label{hda_curv}} \\end{center} \\end{figure} In principle this degeneracy can be broken even with arbitrary $f(z)$ by simultaneous measurements of both Hubble and distance; one can explicitly write $\\Omega_k$ in any FLRW model (with no recourse to the Einstein field equations) as \\cite{CCB,de_degen}: \\be \\Omega_k = \\frac{[H(z)D'(z)]^2-H_0^2}{[H_0D(z)]^2}\\,\\label{curv_ok}, \\ee where $D$ is the dimensionless, transverse, comoving distance with $D(z) = (c/H_0) (1+z) d_A(z)\\,.$ This relation gives the value of $\\Omega_k$ {\\em today}, as a function of measurements at {\\em any} redshift $z$. Hence this can be turned into a powerful test of the Copernican Principle \\cite{copernican}. Since $\\Omega_k$ is a single number, the right hand side will have the same value when measured at any redshift if one is in a FLRW background. If it is found to vary with redshift, then we do not live in a FLRW universe.  The beauty of \\bao~is that they provide both $d_A(z)$ and $H(z)$ using almost completely linear physics (unlike SNIa for example which involve highly complex, nonlinear, poorly understood stellar explosions). In addition, they offer the as yet unproven possibility of delivering constraints on growth though the change in the amplitude of the power spectrum. The time-dependence of the matter density perturbations, $\\delta \\rho/\\rho$ obeys the equation \\begin{equation} \\ddot{\\delta} + 2H\\dot{\\delta} = 4\\pi G \\rho_m \\delta. \\label{deltaeq} \\end{equation} The time-dependence of the growing mode of this equation is given by the growth function, $G(z)$ which in a flat, $\\Lambda$CDM model satisfies \\cite{growth_eis}: \\begin{equation} G(z) = \\frac{5 \\om E(z)}{2} \\intzinf \\frac{(1+z')dz'}{E(z')^3}\\,, \\label{growthz} \\end{equation} while in a general universe with curvature and dark energy dynamics Eq.~(\\ref{deltaeq}) can be rewritten as \\cite{fisher_release}: \\be \\delta'' + \\frac{3}{2}\\left(1 +\\frac{\\Omega_k(x)}{3} - w(x)\\Omega_{\\mathrm{DE}}(x)\\right)\\frac{\\delta'}{x} -\\frac{3}{2}\\Omega_m(x)\\frac{\\delta}{x^2} = 0, \\label{newg} \\ee in terms of the dimensionless scale factor $x = a/a_0 = 1/(1+z),$ where $a_0 = c/H_0(\\sqrt{\\Omega_k})$ is the radius of curvature of the universe. From Eq.~(\\ref{growthz}) we can see that the growth contains a mixture of Hubble and ``distance\" information - as a result measurements of growth are potentially powerful probes of dark energy.  \\subsection{Statistical Standard Rulers\\label{stat}} To illustrate the idea underlying Statistical Standard Rulers (SSR), imagine that all galaxies were positioned at the intersections of a regular three-dimensional grid of known spacing $L$. Measuring angular diameter distances as a function of redshift would be trivial in this case (c.f. Eq.~(\\ref{da_eq})) and we would also have the expansion rate as a function of redshift, measured at a discrete set of redshifts corresponding to the mid-points between galaxies. Now imagine that we start to randomly insert galaxies into this regular grid. As the number of randomly distributed galaxies increases the regular grid pattern will rapidly become hard to see by eye. However, the underlying grid pattern would still be detectable statistically, for example in the Fourier transform.  However, a regular grid distributed throughout space would provide an absolute reference frame and would break the continuous homogeneity of space down to a discrete subgroup (formed by those translations which are multiples of the grid spacing $L$). To get to the core of SSR consider the following prescription for building up a galaxy distribution. Throw down a galaxy at random. Now with some fixed probability, $p$, put another galaxy at a distance $L$ (in any direction) from it. Using the new galaxy as starting point repeat this process until you have the desired number of galaxies. Now there is no regular grid of galaxies but $L$ is still a preferred length in the distribution of the galaxies, and forms an SSR. To reconstruct this preferred scale is a statistical problem. This is illustrated schematically in Figure~(\\ref{rings:fig}), which shows many rings of the same characteristic radius $L$ superimposed on one another. This superposition of rings on the plane visually `hides' the characteristic scale, as number of rings increases, and the sampling of each individual ring is reduced.  \\begin{figure*}[htbp!] \\begin{center} $\\begin{array}{@{\\hspace{-0.6in}}c@{\\hspace{-0.3in}}c} \\includegraphics[width = 3.5in,height = 3in]{rings_200_400.eps} & \\includegraphics[width = 3.5in, height = 3in]{rings_875_100.eps} \\\\ [0.0cm] \\end{array}$ \\caption{Rings of power superposed. Schematic galaxy distribution formed by placing the galaxies on rings of the same characteristic radius $L.$ The preferred radial scale is clearly visible in the left hand panel with many galaxies per ring. The right hand panel shows a more realistic scenario - with many rings and relatively few galaxies per ring, implying that the preferred scale can only be recovered statistically. \\label{rings:fig}} \\end{center} \\end{figure*}  BAO, as we discuss below, provide an elegant SSR hidden between the rest of the galaxy clustering, but they are not the only possible SSR's. Any preferred scale in the clustering provides either an opportunity to apply a relative (absolute) Alcock-Paczynski test \\cite{aptest} in the case the we don't (do) know {\\em a priori} the size of the preferred scale, $L$. Other preferred scales include the Hubble scale at matter-radiation equality (which controls the scale of the turnover in the matter power spectrum) and the Silk damping scale. There may well be other preferred scales imprinted into the primordial clustering of matter. These can be naturally achieved if one inserts a short period of fast rolling into the otherwise slow-roll of inflation. The sudden change of inflaton velocity creates a bump in the matter power spectrum that can serve the same purpose as the BAO. The required fast-roll can be achieved in multi-field models of inflation. \\begin{figure}[htbp!] \\begin{center} \\includegraphics[width = 3.8in, height = 3.1in]{ap_schematic.eps} \\caption{The radial length of an object is given by $c dz/H(z)$ where $dz$ is the difference in redshift between the front and back of the object while the transverse size of the object is $d_A(z) \\theta$ and $\\theta$ is its angular size. If one knows that the object is spherical (but does not know the actual diameter) then one has the Alcock-Paczynski test which gives the product $d_A(z) H(z)$ from measuring $dz/\\theta$. If, as in the case of BAO, one can theoretically determine the diameter, one has the bonus of finding $d_A(z)$ and $H(z)$ separately. \\label{ap_schematic}} \\end{center} \\end{figure}  The SSR provided by the BAO has an additional advantage: it is primarily a linear physics phenomenon, which means we can ignore nonlinear effects\\footnote{These nonlinear effects include redshift space distortions and nonlinear gravitational clustering, which will in general change the spherical nature of the oscillation scale.} to good approximation (we will discuss them later however). This also means we can turn the BAO into a calibrated or absolute Alcock-Paczynski test since the characteristic scale of the BAO is set by the sound horizon at decoupling. As a result the angular diameter distance and Hubble rate can be obtained separately. The characteristic scale, $s_{||}(z)$, along the line-of-sight provides a measurement of the Hubble parameter through \\be H(z) = \\frac{c\\Delta z}{s_{||}(z)}\\label{pll}\\,,\\ee while the tangential mode provides a measurement of the angular diameter distance, \\be d_A(z)  = \\frac{s_{\\perp}}{\\Delta \\theta (1+z)}\\label{perp}\\,.\\ee This is illustrated in the schematic Figure.~(\\ref{ap_schematic}). The horizontal axis is Eq.~(\\ref{pll}) and the vertical axis is $c\\Delta z/H(z),$ Eq.~(\\ref{perp}).   \\begin{figure}[htbp!] \\begin{center} \\includegraphics[width = 5in]{ap_test_twobinsv4.eps} \\caption{Comparing the Alcock-Paczynski (AP) test and BAO constraints -- on the dark energy parameters $w_0, w_a$ in the CPL parameterisation. The solid blue and dashed dark brown lines ellipses illustrate the constraints from a pair of $1\\%$ measurements at $z=1,3$ of the angular diameter distance and Hubble parameter respectively. The filled beige ellipse shows the constraints from a BAO-like survey, where both the Hubble rate and angular diameter distance are measured simultaneously at $z=1,3$ (the errors on each have been increased by $\\sqrt{2}$). In contrast, the dotted line illustrates constraints from the AP test: where $1.5\\%$ measurements are made on the product $H\\times d_A$ at the same redshifts. The AP test provides comparable constraints to using only a single measurement while the BAO provide the tightest constraints; they measure $d_A$ and $H$ simultaneously. The assumed fiducial model is $(H_0, \\Omega_m, \\Omega_k, w_0, w_a) = (70, 0.3, 0, -1,0),$ and the priors on the model are given as $\\mathrm{Prior}(H_0, \\Omega_m, \\Omega_k, w_0, w_a) = (10^4, 10^4, 10^4, 0, 0).$ Figure produced using \\name~\\cite{fisher_release}.  \\label{aptest:fig}} \\end{center} \\end{figure}  While the AP test on its own constrains the product $d_A(z) \\times H(z),$ it is just one function, and so combining the measurements of $d_A$ and $H$ through the BAO provide tighter constraints on cosmological parameters. This is illustrated in Figure~(\\ref{aptest:fig}), which shows the error ellipse in the dark energy parameters from a hypothetical galaxy redshift survey with constraints from the angular diameter distance and Hubble parameter and the product $H \\times d_A$ from the AP test. Constraints on the dark energy parameters from the AP test are similar to those from a single observable such as $d_A$, while the constraints are significantly improved when combining measurements of both $H$ and $d_A.$  One method of extracting a statistical scale from the clustering of galaxies is via the two-point correlation function, $\\xi(r),$ which quantifies the excess clustering on a given scale relative to a uniform distribution with the same mean density. The correlation function of galaxies is approximately described by a power law \\cite{totsuji}, \\be \\xi(r) \\propto \\left(\\frac{r_0}{r}\\right)^{\\gamma}\\,,\\ee with $r_0 \\sim 5 h^{-1}\\mathrm{Mpc}^{-1}.$  A characteristic scale in the clustering of galaxies will appear as a peak or dip in the correlation function, depending on whether there is an excess or deficiency of clustering at that scale. Any characteristic features will also be present in the power spectrum, since the correlation function and power spectrum (we consider for simplicity the simple 1-dimensional spherically averaged power spectrum) form a Fourier pair:  \\be P(k) = \\int_{-\\infty}^{\\infty} \\xi(r)\\exp(-ikr)r^2dr\\,.\\ee We will now see how features in the two functions are related. A $\\delta$ function at a characteristic scale, say $r_*,$ in $\\xi(r)$ will result in power spectrum oscillations, $P(k) \\propto e^{-ikr_*}, $ as can be seen in Figure~(\\ref{schematic:fourier}). These are the \\bao. \\begin{figure}[htbp!] \\begin{center} $\\begin{array}{@{\\hspace{-0.2in}}c@{\\hspace{-0.22in}}c} \\epsfxsize=3in \\epsffile{correl_delta2.eps} & \\epsfxsize=3in \\epsffile{pk_delta2.eps}\\\\ [0.0cm] \\end{array}$ \\caption{Schematic illustration of the Fourier pairs $\\xi(r), P(k)$. A sharp peak in the correlation function (left panel) corresponds to a series of oscillations in $P(k)$ (right panel). The Baryon Acoustic Peak in the correlation function will induce characteristic \\bao~in the power spectrum. \\label{schematic:fourier}} \\end{center} \\end{figure} \\\\ \\begin{figure}[htbp!] \\begin{center} \\includegraphics[width = 3.8in]{pk_bao_omm.ps} \\caption{Changing cosmology moves the \\bao~-- the ``wiggles-only'' power spectra of the SDSS LRG survey. In each panel, a fiducial model with a different value of $\\Omega_m$ has been used to convert the data from redshift to comoving distance, from $\\om = 0.1$ (top) to $\\om = 0.4$ (bottom). The solid lines in each panel show the CDM prediction for the BAO assuming the particular value of $\\om$, while the dashed lines show the same model without the low-redshift small-scale damping term. In each case the baryon fraction is held fixed at $17\\%$, and $h = 0.73$. As the matter density is changed, both the theoretical scale of the BAO changes, and the data moves around through the conversion from redshift to comoving distance. Figure from Percival {\\em et al.,} 2006 \\cite{percival_lrg}. \\label{percival_lrg:fig}} \\end{center} \\end{figure} These characteristic oscillations are powerful probes of dark energy. If we used only the radial Fourier modes we would obtain a measurement of $H(z)$ while the purely transverse modes yield a pure measurement of $d_A(z)$. If we limited ourselves to just the purely radial and transverse modes we would be throwing away a large number of modes (and hence information) since most Fourier modes are not purely radial or transverse of course, but rather contain components of both. As a result, the measurements of $H(z)$ and $d_A(z)$ from actual BAO surveys are actually partially anti-correlated and hence not independent. This anti-correlation actually leads to stronger cosmological constraints than an uncorrelated analysis would suggest, as described in \\cite{seo_eis07}.  A nice example of their sensitivity and the issues involved in using BAO for parameter estimation is evident in Fig.~(\\ref{percival_lrg:fig}) from Percival {\\em et al.} \\cite{percival_lrg} which shows both the SDSS LRG data and theoretical fits as a function of $\\Omega_m$. The panels each show the residual BAO oscillations relative to a non-oscillatory reference spectrum (a technique advanced in Blake {\\em et al.} \\cite{blake_glaze_simulation}), thus isolating the BAO information alone. The panels have a fixed baryon fraction of $17\\%$ relative to dark matter. An important point to take away from this plot is that both the theory {\\em and the data} vary with $\\Omega_m$.  This situation is different from more standard model fitting exercises where the data is static as the model parameters are varied. Here, to compute the data points, a redshift-distance relation (which depends on $\\Omega_m$ of course) must be assumed to move the points into physical, as opposed to redshift space, after which the power spectrum is constructed. In \\cite{percival_lrg}, 31 data sets and theoretical predictions were computed for various flat $\\Lambda$CDM models with different values of $\\Omega_m$. The impact of varying $\\Omega_m$ is seen mostly clearly on the theoretical curves where the amplitude and frequency of oscillation is affected, both by the change to the underlying BAO scale (since the $\\rho_b/\\rho_{\\gamma}$ ratio changes and hence the speed of sound prior to decoupling) and the redshift-distance relation which will move galaxies in the radial direction, bringing them either closer or moving them further away in distance. This dependence of the data on the assumed cosmology is typically ignored in parameter estimation studies using the BAO results, an approximation that is fairly good if one does not range far from the fiducial model \\cite{tegmark_lrg, percival_0705.3323}. The dilation resulting from using a different fiducial model would scale the $k$-axis of the dimensionless power spectrum by the ratio \\cite{eisenstein_05,tegmark_lrg}: \\begin{equation} a = \\frac{d_V(z)}{d_V^{\\mathrm{fiducial}}(z)}, \\end{equation} where $d_V(z) \\equiv \\( d_A(z)^2cz/H(z)\\)^{1/3} $ combines the radial and transverse dilation.  Tegmark {\\em et al.} calibrate this for the parameter range allowed by the combined SDSS-LRG and WMAP3 data, and find that corrections to the $k$-scale are  typically less than or equal to $3\\%.$ As an example, using a fiducial model with $\\om = 0.25$ results in a bias of the measured value of $\\om$ of $\\sim2\\%.$ \\subsection{Physics of the BAO \\label{linphys}} \\begin{figure}[htbp!] \\centerline{\\epsfxsize=2.9in\\epsffile{onez.010.eps} \\quad   \\epsfxsize=2.9in\\epsffile{onez.040.eps}} \\centerline{\\epsfxsize=2.9in\\epsffile{onez.060.eps} \\quad   \\epsfxsize=2.9in\\epsffile{onez.090.eps}} \\centerline{\\epsfxsize=2.9in\\epsffile{onez.250.eps} \\quad   \\epsfxsize=2.9in\\epsffile{onez.259.eps}} \\caption{{\\bf Snapshots of an evolving spherical density perturbation --} the radial mass profile as a function of comoving radius for an initially point-like overdensity located at the origin. The perturbations in the dark matter (black curve), baryons (blue curve), photons (red) and neutrinos (green) evolve from early times ($z=6824$, top left) to long after decoupling ($z=10$, bottom right). Initially the density perturbation propagates through the photons and baryons as a single pulse (top left-hand panel). The drag of the photons and baryons on the dark matter is visible in the top right panel; the dark matter only interacts gravitationally and therefore its perturbation lags behind that of the tightly coupled plasma. During recombination, however, the photons start to ``leak'' away from the baryons (middle left panel); and once recombination is complete $(z = 470$, middle right) the photons freely steam away leaving only a density perturbation in the baryons around $150\\mathrm{Mpc},$ and a dark matter perturbation near the origin. In the bottom two panels we see the how the gravitational interaction between the dark matter and the baryons affects the peak: dark matter pulls the baryons to the peak in the density near zero radius, while the baryons continue to drag the dark matter overdensity towards the $150\\mathrm{Mpc}$ peak (bottom left), finally yielding a peak in the radial mass profile of the dark matter at the scale set by the distance the baryon-photon acoustic wave could have travelled in the time before recoupling. Figure taken from Eisenstein {\\em et al.} \\cite{ESW}. \\label{eisenstein_shell:fig}} \\end{figure}  Before recombination and decoupling the universe consisted of a hot plasma of photons and baryons which were tightly coupled via Thomson scattering. The competing forces of radiation pressure and gravity set up oscillations in the photon fluid. If we consider a single, spherical density perturbation in the tightly coupled baryon-photon plasma it will propagate outwards as an acoustic wave with a speed $c_s = c/\\sqrt{3(1+R)},$ where $R \\equiv 3\\rho_b/4\\rho_\\gamma \\propto \\Omega_b/(1+z)$ \\cite{ESW}. At recombination the cosmos becomes neutral and the pressure on the baryons is removed. The baryon wave stalls while the photons freely propagate away forming what we now observe as the Cosmic Microwave Background (CMB). The characteristic radius of the spherical shell formed when the baryon wave stalled is imprinted on the distribution of the baryons as a density excess. The baryons and dark matter interact though gravity, and so the dark matter also preferentially clumps on this scale. There is thus a increased probability that a galaxy will form somewhere in the higher density remains of the stalled baryon wave than either side of the shell. This scenario is illustrated in the iconic Figure~(\\ref{eisenstein_shell:fig}) from Eisenstein {\\em et al.} \\cite{ESW}.   Had a galaxy formed at the centre of our initial density perturbation there would be a bump in the two-point correlation function at the radius $s$ of our spherical shell, reflecting the higher probability of finding two galaxies separated by the distance $s$. The scale $s$ is usually close to the sound horizon\\footnote{This is perhaps a slight misnomer since $s$ is not the maximum distance that could have been travelled by any fluid. A simple scalar field has speed of sound equal to $c$ and hence travels faster than the baryon-photon sound wave we are discussing here. }, the comoving distance a sound wave could have travelled in the photon-baryon fluid by the time of decoupling, and depends on the baryon and matter densities via \\cite{eis_white}: \\be s = \\int_{z_{rec}}^\\infty\\frac{c_sdz}{H(z)} = \\frac{1}{\\sqrt{\\Omega_mH_0^2}}\\frac{2c}{\\sqrt{3z_{eq}R_{eq}}}\\ln \\left[\\frac{\\sqrt{1+R_{rec}}+ \\sqrt{R_{rec}+R_{eq}}}{1+\\sqrt{R_{eq}}}\\right]\\,, \\ee where $R \\equiv 3\\rho_b/4\\rho_\\gamma \\propto \\Omega_bh^2/(1+z), z_{eq} = \\Omega_m/\\Omega_{rad}$ is the redshift of matter-radiation equality and ``rec'' refers to recombination. The CMB strongly constrains the matter and baryon densities at decoupling and hence the sound horizon, $146.8 \\pm1.8 \\mathrm{Mpc}$ \\cite{wmap5}\\footnote{Confusion sometimes arises from quoting the BAO scale as both $\\sim 150 \\mathrm{Mpc}$ and $105 h^{-1}\\mathrm{Mpc}$. Here $h\\simeq 0.7$ is the Hubble constant today in units of $100\\mathrm{km/s/Mpc}$ and provides the required conversion factor.}. Hence this scale is itself an excellent standard ruler as long as one can measure $\\Omega_b$ to high precision, and in the case where one allows for exotic radiation components in the early Universe, the redshift of equality \\cite{eis_white}. Of course, the early universe was permeated by many such spherical acoustic waves and hence the final density distribution is a linear superposition of the small-amplitude sound waves in the usual Green's function sense. Indeed this real space approach to cosmic perturbations based on Green's functions can be made rigorous, see Bashinsky and Bertschinger \\cite{bert}. ",
        "conclusions": "In the era of precision cosmology, standard rulers of ever-increasing accuracy will provide powerful constraints on dark energy and other cosmic parameters. The Baryon Acoustic Oscillations are rooted primarily in linear physics with nonlinearities that can be well-modelled and corrected for. As a result the characteristic scale of these `frozen relics' imprinted into the cosmic plasma before decoupling will likely remain as the most reliable of the Statistical Standard Rulers in the coming decade."
    },
    "0910/0910.0415_arXiv.txt": {
        "abstract": "We developed an iterative technique to better characterize stellar populations and the central activity of barred  galaxies using evolutionary synthesis codes and OASIS data. The case of NGC\\,5430 is presented here. Our results are reinforcing the role played by the bar and nuclear structures for the evolution of galaxies. ",
        "introduction": "The assemblage of the components of a galaxy is a puzzling but important subject.  Among the key elements, a large scale bar is believed to be  an efficient way to drive gas toward the central kiloparsecs of the galaxy (\\cite{norman}). This may trigger star formation or turn on the central engine (\\cite{haan}). Interactions, mergers, or cosmic filaments are considered to play a role in the accretion of gas and its flow process as well (\\cite{combes}). One way to deepen our understanding of these phenomena is to study in details the content of galaxies.  We used the spectro-imager OASIS at the CFHT and WHT to describe the stellar populations in the central region of 8 galaxies. We developed an iterative technique to separate the flux from two stellar populations, a young population ($<$15\\,Myr) and an old underlying one (see Cantin et al. 2009 for more details).  In summary, the young population is characterized  by comparing the equivalent width of H$\\alpha$ and H$\\beta$ with starburst models from LavalSB (\\cite{dionne}). The old population is studied using absorption lines (e.g. Mg$_2$, Fe{\\sc{i}}) and models of \\cite{gonz}. The iterative process takes into account the flux and line indicators from one population while studying the other one. It allows us, for example, to isolate the absorption component in H$\\beta$  and therefore get more accurate  age and mass estimates. Reliable uncertainties,  within the models considered, are obtained using Bayesian statistics.  \\firstsection ",
        "conclusions": "For most galaxies in our sample, we find:  1) star forming knots of age 5-10 Myr distributed in a nuclear bar, spiral, or ring; 2) a peculiar dust distribution; 3) an underlying old (0.1-5 Gyr) generation of stars with a metallicity lower than the gas; and 4) in most cases, a non thermal activity, either from the composite or the LINER type and not necessarily restricted to the nucleus. We  see relations between the star forming knots morphology, the dust distribution, the variation in the gas and stellar populations abundance, along with the galaxy large scale bar orientation. These relations confirm the importance of the bar for the galaxy evolution. Furthermore in late type spirals, secular evolution, as defined by \\cite{kk},  has been proposed to rearrange the material from the disk and build up a pseudobulge. It has been described as a slow process, either internal  or environmental. This could be a simple scenario, involving here the galaxy bar, to explain the different stellar populations seen in our study.   \\firstsection"
    },
    "0910/0910.5837_arXiv.txt": {
        "abstract": "It has generally been thought that in perturbed Friedmann-Lema\\^{\\i}tre-Robertson-Walker models of the Universe, global topology should not have any feedback effects on dynamics.  However, a weak-field limit heuristical argument, assuming a finite particle horizon for the transmission of gravitational signals, shows that a residual acceleration effect can occur.  The nature of this effect differs algebraically between different constant curvature 3-manifolds.  This potentially provides a selection mechanism for the 3-manifold of comoving space. ",
        "introduction": "The geometrical and global topological freedom allowed by constant curvature Riemannian 3-manifolds as a model of the comoving spatial section of a relativistic universe \\cite{deSitt17,Fried23,Fried24,Lemaitre31ell,Rob35} imply that a Friedmann-Lema\\^{\\i}tre-Robertson-Walker (FLRW) model may be finite in total comoving spatial volume for any of the three curvatures: negative, zero or positive. This resolves the historical dilemma in cosmology according to which two physically desirable qualities of space seemed to be contradictory: a finite universe without boundaries seemed to be a self-contradiction. Riemannian geometry gave a solid mathematical foundation to the resolution of this apparent conflict. Even before the relavistic era, Riemannian 3-manifolds of curvature and/or topology different to that of simply-connected Euclidean space were proposed as models of space by Karl Schwarzschild \\cite{Schw00,Schw98}.  Skipping forward, the start of the ``modern'' epoch of active cosmic topology research can probably be dated back to the two 1993 papers by Starobinsky \\cite{Star93} and Stevens et al. \\cite{Stevens93}.  It is a great pleasure to have Alexei Starobinsky here with us in this meeting.  For reviews of the historical and modern theoretical and observational aspects of cosmic topology, see refs \\cite{LaLu95,Lum98,Stark98,LR99,BR99,RG04}.   Although, in general, these 3-manifolds are not vector spaces, Grassmann's pioneering work in linear algebra \\cite{Grass44A1,Grass62A2} nevertheless provides useful tools (as it does in nearly all of modern science), since working in the covering spaces $ \\mathbb{H}^3 $, $\\mathbb{R}^3$, and $ S^3 $, for negative, zero, and positive curvature, respectively, is useful for many observational and theoretical purposes.  The covering space $\\mathbb{R}^3$ is a vector space, and calculations made when the covering space is the hypersphere $S^3$ are very straightforward to program when $S^3$ is embedded in $\\mathbb{R}^4$. Both the analytical and numerical calculations referred to below for the spherical well-proportioned spaces specifically use this latter technique---thanks to Hermann Grassmann.  What was long taken as self-evident in cosmic topology research was the inference that since the Einstein field equations are local, there is no way that the global topology of the spatial section of an FLRW universe could have an effect on that universe's dynamics. The only known link between topology and dynamics was through curvature, since the three different curvatures allow three different families of 3-manifolds. ",
        "conclusions": ""
    },
    "0910/0910.0848_arXiv.txt": {
        "abstract": "Solar gravity modes (or g modes) -- oscillations of the solar interior for which buoyancy acts as the restoring force -- have the potential to provide unprecedented inference on the structure and dynamics of the solar core, inference that is not possible with the well observed acoustic modes (or p modes). The high amplitude of the g-mode eigenfunctions in the core and the evanesence of the modes in the convection zone make the modes particularly sensitive to the physical and dynamical conditions in the core. Owing to the existence of the convection zone, the g modes have very low amplitudes at photospheric levels, which makes the modes extremely hard to detect. In this paper, we review the current state of play regarding attempts to detect g modes. We review the theory of g modes, including theoretical estimation of the g-mode frequencies, amplitudes and damping rates. Then we go on to discuss the techniques that have been used to try to detect g modes. We review results in the literature, and finish by looking to the future, and the potential advances that can be made -- from both data and data-analysis perspectives -- to give unambiguous detections of individual g modes. The review ends by concluding that, at the time of writing, there is indeed a consensus amongst the authors that {\\it there is currently no undisputed detection of solar g  modes.} ",
        "introduction": "\\label{intro}  The detection of solar oscillations by \\citet{Leighton1962} marked the start of a memorable era for solar physics, and the birth of a new field, helioseismology.  The oscillatory motions detected in the solar atmosphere had typical periods of about five minutes, and \\citet{Deubner75} subsequently verified observationally that the oscillations were the visible manifestation of acoustic (p) modes trapped in cavities in the solar interior. Later in the 1970s, \\citet{AC79} showed that the Sun supported truly global modes of oscillation. The low-angular-degree (low-$l$) modes observed by \\citet{AC79} penetrated the solar core, and had therefore opened a window on the structure of the deep interior of the Sun.  \\citet{GG80} subsequently made long-duration observations of the Sun from the South Pole, providing the very first clean power spectrum free from day-night interruptions.  While inference on the internal structure of the Sun made using high-quality data on the p modes has gone from strength to strength over the intervening 30 years, attempts to detect g modes have proven altogether more elusive.  This is in stark contrast with the ubiquitous detection of stellar g modes across the Hertzsprung-Russell diagram such as in white dwarfs \\citep{Winget2008}, in subdwarfs B stars \\citep{Green2003}, in B stars \\citep{Waelkens1991} and in $\\gamma$ Dor stars \\citep{Aerts1998}.  The detection, and subsequent exploitation, of data on the solar g modes would provide much more precise inference on the structure and dynamics of the solar core than is possible with p modes. This is because the g modes are trapped in cavities in the radiative interior, and they may have much higher displacement amplitude in the radiative zone / core than their acoustic counterparts. However, this high sensitivity to the core properties comes at a price. Because the g modes are trapped beneath the convective envelope they have very low amplitudes at the photospheric level where the perturbations due to the oscillations are detected; hence their elusive nature.  The first attempt to detect long-period oscillations date back to the mid 1970s, e.g., \\citet{Severnyi1976}. Further claims, from ground-based Doppler velocity observations, were made in the 1980s, e.g., \\citet{PDPS83}. It was realized very early on in the development of helioseismology as a field that a space mission would be ideally suited not only for observing the p modes but also for detecting g modes (in principle providing the long-term stability needed to detect the low-frequency g modes).  The DISCO\\footnote{Dual Spectral Irradiance and Solar Constant Orbiter} mission was proposed by \\citet{Bonnet81} primarily as an irradiance mission to which helioseismology was added later on. The initial design was subsequently modified to accommodate several in-situ instruments and spectrographs, and finally matured into the SOHO\\footnote{SOlar and Heliospheric Observatory} mission \\citep{Domingo1995}.  One of the goals of the helioseismology instruments on SOHO was to detect g modes. Following its launch in 1995 December, initial analyses of the helioseismology data failed to uncover evidence for the modes. By 1997, it was realized that the searches would benefit greatly from an internationally coordinated effort comprised of the principal observational and theoretical scientists involved in g-mode research. This led to the formation of the Phoebus Group, which has been working since to use all available high-quality helioseismic data, and to develop the analysis techniques that are required, to detect g modes.  In this review we collect and summarize the knowledge accumulated by the Phoebus group over more than a decade of g-mode research.  Our review is broken down as follows. In Section~2 we review the theoretical properties of the solar oscillations. [A complete presentation of the properties of solar and stellar oscillations may be found for example in \\citet{Unno89}, \\citet{jcdgb91}, and \\citet{2002RvMP...74.1073C}.]  We also consider the accuracy of the theoretical computations of the g mode frequencies, and detail the various factors that contribute to the uncertainties (i.e., the accuracy of the oscillation computations themselves, and the sensitivity of the frequencies to the structure and dynamics of the solar models).  In Section~3 we discuss the excitation and damping of the g modes, again from a theoretical perspective. We consider the different physical mechanisms that can potentially contribute to the stability of g modes, and present theoretical estimates of the amplitudes of g modes excited stochastically by convection in the convective envelope.  In Section~4, we turn to the observations. We begin by considering the different detection techniques that are available to observers, and discuss in detail the constituent elements required to ``build'' a front-to-end analysis algorithm. Our presentation quite naturally includes an in-depth discussion of frequency-domain analysis, and hypothesis testing (frequentist versus Bayesian). We finish this section by reviewing (from the literature) attempts to apply these techniques to detect g modes.  Finally, in Section~5, we conclude with some remarks on the current state of play. Our view is that, at the time of writing, there is no unambiguous evidence for the detection of solar g modes, although some of the most recent theoretical predictions suggest a positive result may be attainable in the near future.  We therefore finish by looking to the future prospects for finally detecting g modes. ",
        "conclusions": "Figure~\\ref{fig:amp_dog_kumar_kevin} shows the comparison of the best observational measurements with theoretical amplitudes of the g modes. The theoretical g-mode amplitudes range from 10$^{-2}$\\,mm\\,s$^{-1}$ to a few mm\\,s$^{-1}$.  The most optimistic theoretical mode amplitudes could be collectively detectable, after the collection of more than 10 years of observations of the solar radial velocity.  Unfortunately, the detection of individual peaks is very far from being feasible since the only possibility would be to have the most optimistic amplitude larger by 50\\%, or twice as much in power \\citep{KB2009}.  Since the noise reduction scales with the observing time $T$ like $\\log(T)/\\sqrt{T}$, another 80 years or so of data would be required to reduce the actual detection limit by a factor two.  The limit does not scale like $1/\\sqrt{T}$ because the probability limit is kept constant in a given detection window while the number of frequency bins in that window increases like $T$ \\citep{TA99}. At the time of writing, there is indeed a consensus amongst the authors of this review that {\\it there is currently no undisputed detection of solar g  modes.}   Despite this state of affairs, the search is not over yet.  What can be done to turn our fortunes around? First of all, some of the detection techniques presented in this review paper that rely on the theoretical knowledge presented in Sections 2 and 3. Several are yet to exploit fully the theoretical information: this calls for the application of full Bayesian statistical inference on the data (yet to be done).  The observation of signals from the the solar limb will hopefully offer another way of improving the potential visibility of the g modes.  These observations will be carried out by PICARD, which is due to be launched in 2009.  Because the solar atmosphere is relatively noisy, reductions of noise levels can potentially be achieved by making multiple observations that are sensitive to perturbations at different heights in the atmosphere. Here, one would seek to rely on changes in the coherence of the noise with height.  The GOLF-NG instrument recently put in operation uses this principle for measuring solar radial velocities in 8 different heights of the solar atmosphere \\citep{STC2006,Salabert2008b}.  Development of methods based on time-distance helioseismology offer promises not only for detecting directly the g modes, but also for probing the solar core by making observations from two widely separated point of view.  The technique is called stereoscopic seismology, and may provide the structure and the dynamics of the core with a completely different approach with the Solar Orbiter mission \\citep{AG2003,Gizon2006}.  Finally, we could try to detect the perturbation of the gravitational potential created by the g modes.  For instance, the LISA\\footnote{The ESA/NASA gravitational wave laser interferometric space antenna} mission may provide a lower detection limit than the current classical helioseismic techniques.  \\citet{Polnarev2009} deduced a lower detection limit for LISA that is about a factor five lower in amplitude than the actual GOLF limit for modes of $m=\\pm2$ (see Figure~\\ref{fig:amp_dog_kumar_kevin}).  This limit would be low enough to detect g modes if they had amplitudes as high as those predicted by \\citet{KB2009}.  The variations in the potential can be also detected by using laser ranging, e.g., the ASTROD mission \\citep{Ni2007, Appourchaux2009}.  As shown by \\citet{RB2008}, the capability of ASTROD would be sufficient to allow the detection of g modes even if their amplitudes were to be as low as those predicted by \\citet{Kumar96}.  We conclude with an optimistic comment. It is superficially unfortunate that we may not yet have detected g modes in the Sun.  On the brighter side, on which we always prefer to be, it shows that we have before us a greater challenge which will yield greater satisfaction when we overcome it.  We are all now much more prepared to continue the search.  \\acknowledgement The whole Phoebus group would like to thank the long-time support provided by the European Space Agency, in particular we thank Clare Bingham, Cecilia Nillson, Mylene Riemens, Birgit Schroeder for secretarial support, and Martin Huber, Peter Wenzel and Bernard Foing for political and financial support; and the support provided by the International Space Science Institute, in particular we thank Brigitte Schutte, Saliba Saliba, Vittorio Manno and Roger-Maurice Bonnet for supporting our program.  SOHO is a mission of international collaboration between ESA and NASA.  The Phoebus group would like also to thank Boris Dintrans for setting the CVS server at the Observatoire Midi-Pyr\\'en\\'ees which allows for the parallel writing of the review.  \\appendix"
    },
    "0910/0910.0529_arXiv.txt": {
        "abstract": "{The massive post-Main Sequence star W243 in the galactic starburst cluster Westerlund 1 has undergone a spectral transformation from a B2Ia supergiant devoid of emission features in 1981 to an A-type supergiant with a rich emission-line spectrum by 2002/03. The star was proposed as a Luminous Blue Variable undergoing an eruptive event by Clark \\& Negueruela (2004).} {We examine the continued evolution of W243 from 2002 to 2009 to understand its evolutionary state, current physical properties and the origin of its peculiar emission line spectrum.} {We used VLT/UVES and VLT/FLAMES to obtain high resolution, high signal-to-noise spectra on six epochs in 2003/04 (UVES) and ten epochs in 2008/09 (FLAMES). These spectra are used alongside other lower-resolution VLT/FLAMES and NTT/EMMI spectra to follow the evolution of W243 from 2002 to 2009. Non-LTE models are used to determine the physical properties of W243.} {W243 displays a complex, time-varying spectrum with emission lines of Hydrogen, Helium and Lyman-$\\alpha$ pumped metals, forbidden lines of Nitrogen and Iron, and a large number of absorption lines from neutral and singly-ionized metals. Many lines are complex emission/absorption blends, with significant spectral evolution occurring on timescales of just a few days. The LBV has a temperature of $\\sim$8500K (spectral type A3Ia$^{+}$), and displays signs of photospheric pulsations and weak episodic mass loss. Nitrogen is highly overabundant, with Carbon and Oxygen depleted, indicative of surface CNO-processed material and considerable previous mass-loss, although current time-averaged mass-loss rates are low. The emission-line spectrum forms at large radii, when material lost by the LBV in a previous mass-loss event is ionized by an unseen hot companion. Monitoring of the near-infrared spectrum suggests that the star has not changed significantly since it finished evolving to the cool state, close to the Humphreys-Davidson limit, in early 2003.} {} ",
        "introduction": "Luminous Blue Variables (LBVs; \\citealt{hd94}) represent a short-lived transitional stage in the life of the most massive stars as they leave the main sequence and evolve towards the Wolf-Rayet phase. LBVs are characterized by very high rates of mass-loss, dense, slow winds and variability in luminosity and temperature that range from short-term microvariations of $\\sim$0.1 to 0.2 magnitudes to the rare, catastrophic \\object{$\\eta$~Carinae} type eruptions for which the LBV class is most well known. In its quiescent state a LBV appears as a very luminous ($\\sim$10$^6 \\text{L}_{\\odot}$) B-type supergiant that lies on the \\object{S~Doradus} instability strip \\citep{svk04}, but on timescales of a few decades the LBV will move redwards on the Hertzsprung-Russell Diagram (HRD), taking on the appearance of an A- or F-type supergiant and brightening by $m_{v}$$\\sim$1~to~2 due to the change in bolometric correction. More rarely, LBVs may undergo \\textit{giant} eruptions in which bolometric luminosity is not conserved and substantial mass-loss occurs, with many LBVs surrounded by expanding nebule (e.g. \\citealt{w03}) that consist of material ejected from the star during earlier mass-loss events. Finally, there is growing evidence that LBVs may be the progenitors of some type~IIn supernovae (e.g. \\object{SN2006gy}, \\citealt{setal07}; \\object{SN~2005gl}, \\citealt{galyam}), while the bipolar nebula surrounding the candidate-LBV \\object{HD 168625} has a morphology similar to \\object{Sk -69$^{\\circ}$202}, the progenitor of SN1987A \\citep{s07}. However, the rarity of LBVs means that the nature of these stars is still poorly understood, and the mechanisms behind the periodic increases in mass loss as the star leaves the quiescent state and the causes of the giant $\\eta$-Carinae type eruptions are unknown.  A recent addition to the catalogue of galactic LBVs is \\object{W243}\\footnote{RA=16~47~07.5~$\\delta$=-45~52~28.5,~J2000.} \\citep[hereafter Paper~I]{cn04} in the starburst cluster \\object{Westerlund~1} (hereafter Wd~1; \\citealt{w61, cncg05}). Early observations included \\emph{star~G} (=\\object{W243}) amongst a group of OB supergiants, with \\cite{bks70} reporting $m_{v}$$\\sim$15.1$\\pm$0.2\\footnote{We note that photometry from different studies of Wd1 tends to show systematic differences on the order of a few tenths of a magnitude \\citep{clark09} and the values for \\object{W243} listed here may not be directly comparable.} and spectral type B0.5Ia, while \\cite{lock74} list $m_{v}$$\\sim$15.6; see also \\cite{k77}. \\cite{w87} report $m_{v}$$\\sim$14.38 from photometry obtained in 1960--66, but used a spectrum taken in 1981 to classify \\object{W243} as a B2Ia supergiant with a spectrum devoid of emission lines. However, Paper~I found that by 2002/03 \\object{W243} was displaying the spectrum of an A2Ia supergiant, implying an apparent decrease in temperature of $\\sim$10$^4\\text{K}$ during the intervening period, accompanied by the development of a rich emission line spectrum showing lines of H, He~I, and Ca~II. Excellent agreement was found between the spectral types of six other B- and A-type supergiants observed by both \\cite{w87} and \\cite{cncg05}, and Paper~I concluded that \\object{W243} is an LBV that was quiescent when observed by \\cite{w87} but has since undergone a significant outburst, moving it redwards on the H-R diagram and (assuming constant bolometric luminosity) close to the Humphreys-Davidson limit.  More recently, \\object{W243} has been listed as an aperiodic variable with $m_{v}$$\\sim$15.73 by \\cite{bonanos07}, while \\cite{groh06, groh07} report infrared spectra with Pa$\\gamma$ emission that resembles the cool phase of the LBV \\object{HR~Carinae} \\citep{m02} and a K-band spectral morphology that is very similar to other LBVs such as \\object{AG~Carinae} \\citep{groh09}, \\object{LBV~1806-20} \\citep{eik04} and the \\object{Pistol~Star} \\citep{figer98}. However, the strong He~I emission seen at $6678\\text{\\AA}$, $7065\\text{\\AA}$ (Paper~I) and $10830\\text{\\AA}$  \\citep{groh07} is anomalous, and \\cite{groh06} note that the K-band spectrum also implies a higher temperature than that of a typical yellow hypergiant (YHG) and suggest that \\object{W243} may be evolving back towards a hotter state. \\cite{clark08} report \\object{W243} as a weak (L$_{x}<10^{32}$erg s$^{-1}$) X-ray source\\footnote{This is consistent with X-ray observations of other LBVs (e.g. \\citealt{muno06}), although some LBVs (e.g. \\object{$\\eta$~Carinae}, \\object{HD~5980}) have significantly higher X-ray luminosities (\\citealt{muno06}; \\citealt{clark08})}, while \\cite{dougherty} find a current mass-loss rate for \\object{W243} of $4-6\\times10^{-6} {\\rm  M}_\\odot\\,{\\rm yr}^{-1}$ from radio observations. No evidence is seen of the nebulosity typically surrounding LBVs (e.g. \\citealt{w03}), although it is likely that such a nebula would be rapidly disrupted in the environment of Wd~1. Finally, \\cite{ritchie09} report evidence for photospheric pulsations in \\object{W243}; this is examined further in Section~\\ref{sec:pulsations}.  This paper presents further analysis of the Paper~I dataset, along with analysis of five subsequent high-resolution Echelle spectra taken during 2004 and intermediate- and high-resolution spectra from 2005, 2008 and 2009 to examine the physical state and evolution of \\object{W243} over a seven-year timescale. Details of the observations and data reduction are presented in Section~\\ref{sec:obs_data}, results are presented in Section~\\ref{sec:results} and are discussed in Section~\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  \\subsection{The spectrum of W243}  We summarize the spectrum of \\object{W243} as follows: \\begin{itemize} \\item \\textbf{H, He~I, O~I, Mg~I, Ca~II, Fe~II, [N~II] and [Fe~II]} are seen in emission. Strong H$\\alpha$, H$\\beta$ and Pa$\\delta$ emission is seen, with the H$\\alpha$ and H$\\beta$ lines displaying double-peaked profiles, while the higher Paschen-series lines show core emission on the blue absorption wing. He~I emission is highly variable, while O~I $\\lambda$8446 and many Fe~II lines also show significant variability. [N~II] and a few [Fe~II] lines are observed. Apart from a few highly-distorted lines, all permitted and forbidden emission lines display a blueshift of $\\sim$50kms$^{-1}$ with negligible epoch-to-epoch variability.  \\item \\textbf{Many neutral or singly-ionized metals are seen in absorption}. Strong absorption lines of Si~II, N~I and O~I are seen, along with many Fe~II lines; Fe~I and Fe~III features are not observed. Many weaker lines from singly-ionized iron-group elements are also seen. Modeling of the absorption line spectrum shows 8300$\\le$T$_{\\text{eff}}\\le$9300K (A2--A4Ia$^{+}$, with $T_{\\text{eff}}\\sim$8500K preferred). This is strongly inconsistent with the observed He~I emission and Ly$\\alpha$ and Ly$\\beta$ fluorescence lines of O~I and Fe~II, implying the presence of an unseen source of ionizing photons.  \\item \\textbf{Photospheric pulsations are apparent in the metal absorption lines}, with periodic broadening and the development of excess blue-wing absorption suggestive of periods of increased mass loss. However, no evidence is seen for eruptive mass loss, and \\object{W243} appears to be in a quiescent cool-phase state. \\end{itemize}  \\subsection{The current state of W243}  With $L_{*}$=7.3$\\times$10$^{5}L_{\\odot}$ (assuming $d$=4.5kpc; \\citealt{cncg05}) and $T_{\\text{eff}}$$\\sim$8500K, \\object{W243} lies near the post-Main Sequence track of a $\\sim$40$M_{\\odot}$ star (e.g. \\citealt{schaller92}), falling as expected between the $\\sim$40-55$M_{\\odot}$ progenitors of the Wd1 Wolf-Rayet population \\citep{crowther06} and the $\\le$30$M_{\\odot}$ $\\sim$O7-8V stars at the main sequence turn-off \\citep{negueruela09}. \\object{W243} appeared on the blue side of the HRD as a luminous, early-B supergiant when observed by \\cite{bks70}, \\cite{lock74} and \\cite{w87} before undergoing an (unobserved) event that led to \\object{W243} moving redwards to the `yellow void' and taking on its current appearance as an early-A hypergiant. The final stages of this evolution were seen in the early observations described in Paper~I, with the star remaining generally quiescent in its cool-phase state since 2003. Our modeling shows Nitrogen to be highly overabundant, with the depletion of Carbon and Oxygen suggesting we are observing significant quantities of CNO-processed material. Other $\\alpha$-process elements (Ca, Mg, Si) are also enhanced. The implication is therefore that \\object{W243} is either in an advanced pre-RSG LBV phase, having already ejected its less processed, Hydrogen-rich outer layers, or has evolved through the RSG phase and returned to the blue side of the HRD, with the current CNO-rich material dredged up during the RSG phase. \\object{W243} does not display the extended circumstellar ejecta expected for \\textit{either} evolutionary stage, but, as commented on in Paper~I, this is likely to be rapidly dispersed in the extreme environment of Wd1. While the N/O ratio in \\object{W243} is similar to \\object{AG~Car} \\citep{groh09}, the latter object is both more massive and (considerably) more luminous than \\object{W243} and the same evolutionary phase cannot be assumed, while post-RSG objects such as \\object{IRC~+10~420} are also overabundant in Nitrogen \\citep{oudmaijer}\\footnote{Sodium abundances are used to infer a post-RSG phase for \\object{$\\rho$~Cas} (e.g. \\citealt{deJager97}), but the only features apparent in our spectra result from interstellar absorption.}. However, future abundance studies of the unique coeval population of early-B supergiants, YHGs and BIa$^{+}$/WNVL Wolf-Rayet precursors found in Wd1 offer the opportunity to accurately constrain both the evolutionary state of \\object{W243} itself and the general passage of massive stars as they leave the main sequence and evolve through the `zoo' of transitional objects towards the Wolf-Rayet phase.  \\subsection{Comparison with other stars}  \\subsubsection{Cool-phase LBVs}  The visual and near-IR spectrum of \\object{W243} shows many similarities with other cool-phase LBVs. The Si~II doublet is prominent in cool supergiants \\citep{davies05}, and is therefore reported in many objects, while the strong near-IR N~I lines are also seen at $\\sim$B8Ia and later, e.g. in the spectra of \\object{S~Dor} \\citep{munari09} and \\object{HR~Car}\\footnote{The N~I lines are notably stronger in \\object{W243} than in \\object{HR~Car}, implying a later spectral type \\citep{mt99}.} \\citep{m02}. The R-band spectrum of \\object{W243} shows strong similarities with the A2Ia$^{+}$ LBV \\object{HD~160529} \\citep{stahl03, chentsov03}, although subtle differences are present, with the C~II~$\\lambda\\lambda$6578, 6583 doublet notable in the spectrum of \\object{HD~160529} and Ne~I also reported; neither can be robustly identified in the case of \\object{W243}. The rich Fe~II spectrum is also a typical feature of cool-phase LBVs, with \\object{HR~Car} a canonical example, displaying prominent Fe~II emission/absorption blends in the cool states of 1991 and 1998--2002 that fade as the LBV moves bluewards on the HRD. The transition from P~Cygni to inverse P~Cygni profiles in Fe~II seen in Figure~\\ref{fig:FeIIcompare} is also reported in the cases of \\object{S~Dor} \\citep{wolf90, wolf92}, \\object{R127} \\citep{wolf92, walborn08} and \\object{HR~Car} (\\citealt{n97}; \\citealt{m02}; Crowther, priv. comm. 2009); we return to this issue in Section~\\ref{sec:location}.  However, notable differences also exist between \\object{W243} and other early-A LBVs. The H$\\alpha$ and H$\\beta$ lines seen in the spectrum of \\object{W243} are unusual, appearing strongly in emission but lacking the P~Cygni absorption component generally seen in LBV spectra. The strong He~I and Ly$\\alpha$-pumped Fe~II and O~I emission in \\object{W243} are also clearly discrepant with the cool-phase emission and absorption lines and are more in keeping with early-B/Ofpe hot-phase LBVs (e.g. \\object{AG~Car}; \\citealt{groh09}), although other characteristic hot-phase features such as emission lines of Fe~III, N~II, Si~II and Si~III are absent. The He~I~$\\lambda\\lambda$5876, 6678 lines are seen in absorption in the cool-phase spectra of \\object{HR~Car} \\citep{m02}, and the four B6-A2 (candidate) LBVs described by \\cite{chentsov03}, while He~I lines appear in emission in \\object{R127} only as the Fe~II spectrum fades during the transition to the hot, quiescent state \\citep{walborn08}. \\object{HD~160529} provides a striking example of these differences, with the otherwise-similar A2Ia$^{+}$ spectrum showing He~I in absorption and a strong P~Cygni profile in the H$\\alpha$ and H$\\beta$ lines; our model also suggests that He~I should be visible weakly in absorption, rather than the observed strong emission. Infra-red spectra of the mid-B or later (candidate) LBVs listed by \\cite{groh07} again show \\object{W243} to be atypical, with a strong He~I~$\\lambda$1.083$\\mu$m emission line (of slightly greater strength than the adjacent Pa$\\gamma$ line) contrasting with the remainder of the sample that show the line to be in absorption in two cases (\\object{HD~168625}, \\object{HD~168607}), almost absent in \\object{HR~Car} and showing a weak emission component with strong P~Cygni absorption in \\object{HD~160529}.  \\subsubsection{Yellow Hypergiants}  Our modeling shows that the absorption line spectrum of \\object{W243} is well described by a cool hypergiant close to the Eddington limit, and we observe spectral variability that is similar to the well-studied YHGs \\object{$\\rho$~Cas} \\citep{lobel98} and \\object{HR~8752} \\citep{deJager97}. Within Wd1, \\object{W265} also displays pulsational variability in the near-IR N~I lines, while at radio wavelengths \\object{W243} is very similar to the F2Ia$^{+}$ YHG \\object{W4} \\citep{dougherty}. The lack of Fe~I features indicates that \\object{W243} is hotter than these objects, which show strong neutral metal absorption and a complete absence of He~I, with the latter also not observed in the A-type YHGs \\object{W12a} (A5Ia$^{+}$) and \\object{W16a} (A2Ia$^{+}$). The peculiar mid-A YHG \\object{IRC~+10~420} (\\citealt{oud98}; \\citealt{hump02}) shows the strongest spectral similarities to \\object{W243}, with Fe~II multiplet~40 and ~74 lines in emission and a strong, double-peaked H$\\alpha$ emission line with broad electron-scattering wings; both objects also display double-peaked Ca~II emission. However, the close agreement between the H$\\alpha$ and Ca~II features in velocity space (e.g. Figure 5, \\citealt{hump02}) is not seen in \\object{W243}\\footnote{This may be a result of the blending with the adjacent Paschen-series absorption/emission features which are not seen strongly in \\object{IRC~+10~420}}, suggesting that the origin may not be the same. In addition, the 2009 VLT/FLAMES spectra show that the Ca~II infra-red triplet no longer displays a double-peaked profile, with only a weak `step' in the red flank of the emission line hinting at the previous location of the absorption feature. Lack of contemporaneous R-band spectra means that we cannot tell if the double-peaked H$\\alpha$ profile is also now absent, although we note the clear weakening in the redwards peak in the VLT/UVES spectra and its apparent absence in the lower-resolution VLT/FLAMES LR6 spectra from 2005 (see Figure~\\ref{fig:HaPd}). Therefore, although the double-peaked line profiles seen in \\object{W243} may correspond to the disk (e.g. \\citealt{jones93}) or bipolar outflow \\citep{oud94, davies07} models proposed for \\object{IRC~+10~420}\\footnote{\\cite{hump02} also propose a complex \\textit{rain} model for \\object{IRC~+10~420} in which both inflow and outflow are present in an opaque wind. However, this requires a mass-loss rate $\\sim$10$^{3}$ times greater than currently appears to be the case for \\object{W243}.}, the central absorption may also result from radiative transfer effects in the line-forming region; additional modeling would be required to confirm this.  \\subsection{Origin of the emission line spectrum}  While the spectrum of \\object{W243} shows similarities with both cool, A-type LBVs and A-~and~F-type YHGs, the juxtaposition of a cool hypergiant absorption line spectrum with hot-phase LBV He~I and Lyman-pumped Fe~II and O~I emission lines is an obvious discrepancy. However, our model is unable to reproduce \\textit{any} of the observed emission features while simultaneously preserving the N~I and Si~II absorption lines, and we interpret this as an indication that \\object{W243} is a binary, consisting of the cool-phase LBV and an undetected hot OB (or possibly Wolf-Rayet) companion that is the source of the ionizing flux responsible for the observed radio emission and discrepant emission lines. The binary fraction amongst massive stars in Wd~1 is known to be large \\citep{crowther06, clark08, ritchie09}, and although \\object{W243} lacks the strong non-thermal radio and X-ray emission that provide evidence for colliding winds in a binary system, weak X-ray emission is observed \\citep{clark08}. Despite the implied presence of a hot companion to the LBV, radial velocity measurements are not obviously consistent with binarity unless we are viewing \\object{W243} at a highly favourable angle or the companion is in a long-period, possibly highly-eccentric orbit around the LBV primary as has been suggested for \\object{$\\eta$~Carinae} \\citep{nielsen07}. An alternative configuration may be similar to that observed in the YHG \\object{HR~8752}, where the flux from a B1V companion in a wide orbit is responsible for the formation of the observed [N~II] emission and a compact radio nebula, neither of which could be formed by the cool hypergiant itself \\citep{stickland78}. A similar scenario may also apply to the YHG \\object{W265}, which also displays [N~II] and radio emission \\citep{ritchie09}. In the case of \\object{W243} a hotter companion is probably required to provide the ionizing flux necessary to form the observed He~I and Ly$\\alpha$-pumped emission, both absent in \\object{HR~8752} and \\object{W265}, although we note the absence of higher-excitation emission lines such as N~II, [S~III], [Fe~III] and [Ar~III] that are observed in the sgB[e] binary \\object{W9} \\citep{cncg05, clark08}, which also displays far stronger Ly$\\alpha$-pumped lines than those observed in \\object{W243}.  The lack of observed Fe~I or Fe~III features in the spectrum of \\object{W243} is unusual and implies that almost all of the iron is in the Fe$^{+}$ state in the line formation region. This is reminiscent of the B and D~Weigelt blobs of \\object{$\\eta$~Carinae} \\citep{verner02}, which display spectra dominated by Fe~II and [Fe~II] emission originating from upper levels around 6eV (permitted lines) and 2eV (forbidden lines). This Fe~II spectrum arises from collisional excitation at $10^{6}<n_{e}<10^{7}$cm$^{-3}$ populating the lower Fe~II levels, with strong continuum pumping routes to $\\sim$6eV levels arising from the a$^{4}$D and a$^{2}$G levels (the a$^{2}$G level is the upper level of our observed [Fe~II] multiplet 14F emission) and Ly$\\alpha$ pumping from the a$^{4}$D and a$^{4}$G levels; this also appears a plausible model for the Fe~II emission region of \\object{W243}\\footnote{The similarity between the Fe~II-rich spectra of \\object{W243} and other cool-phase LBVs, e.g. \\object{HR~Car} or \\object{S~Dor}, also suggests similar conditions in the Fe~II line formation regions, although we do not suggest that a hot companion is required in all of these objects; instead, the similarity probably reflects the complex balance of collisional, continuum and Ly$\\alpha$-pumping processes in these objects \\citep{verner02}.}. The significant variability in Fe~II, He~I and O~I emission suggests that the efficiency of the pumping channels varies, possibly due to changes in Lyman-series optical depth or other radiative transfer effects.  \\subsection{Location of the emission line formation region} \\label{sec:location}  The transition from P~Cygni to inverse P~Cygni profiles in the Fe~II lines seen in Figure~\\ref{fig:FeIIcompare} is observed in other LBVs, and is interpreted as an indication of infall during the radial contraction phase of the stellar photosphere during the \\object{S~Dor} variability cycle \\citep{wolf90}. In the case of \\object{W243}, this transition occurs at a time when the strong Si~II and N~I absorption lines develop broad blue wings\\footnote{A similar transformation is correlated with periods of enhanced mass-loss in the YHG \\object{$\\rho$~Cas} \\citep{lobel03}.} but inspection of the Fe~II multiplet~74 and higher Paschen-series lines suggests that these may not be \\textit{true} (inverse) P~Cygni profiles, and rather an effect of emission at near-constant radial velocity superimposed on an absorption line with a varying line centre: this is similar, but more pronounced, to the effects seen in \\object{$\\rho$~Cas} where static emission lines and variable absorption lines produce apparent line-splitting (see \\citealt{lobel98}, figs.~13 and~15). With \\object{W243}, if the absorption is at its most blueshifted (and especially if a broad blue wing also develops) the emission/absorption blend will take on an apparent P~Cygni profile (e.g. as visible in the Pa$\\delta$ line in the left panel of Figure~\\ref{fig:HaPd} in the 10/07/2004, MJD=53192.2 spectrum, or in the Fe~II multiplet~40 lines in Figure~\\ref{fig:FeIIcompare}), while at its most redshifted the absorption component will be seen entirely on the red flank of the emission component, creating an inverse P~Cygni profile: however, neither case corresponds to the significant outflow or inflow normally associated with these spectral features.  The near-static radial velocities of the undistorted emission lines over the course of our observations and the close correspondence between the radial velocities of the permitted and forbidden lines suggests that they all form at significant radii, when material expanding slowly at constant velocity is ionized by the flux from the hot secondary. As an emission spectrum was not observed by \\cite{w87}, it appears that this material is a relic of mass-loss during the transition from an early-B supergiant to the current cool state. The O~I~$\\lambda$8446, Fe~II and higher Paschen-series lines have very similar profiles, and all likely form in a common region in the boundary zone between H~I and H~II regions required for efficient Ly$\\beta$-pumping of the O~I line, while the consistency between the measured radial velocities of almost all the permitted and forbidden emission lines suggests formation in the same comoving region. The radial velocity of the Ca~II emission is somewhat variable, while the line is broader than all but the Balmer-series emission lines, most probably reflecting formation in a different location: the $\\sim$11.9eV ionization potential means recombination will be negligible in a H~II region, and the line likely results from resonance absorption in an H~I region rather than being limited to the H~I/H~II transition zone. It is plausible that the Ca~II emission may be a result of dust evaporation (e.g. \\citealt{prieto}) which would indicate a post-RSG state for \\object{W243}, although the [Ca II] emission that would provide stronger support for this hypothesis is absent \\citep{shields99}. These lines lie in a region of heavy telluric contamination, but we also note that [Ca~II] is only weakly detected in the B[e] star \\object{MWC 349A} despite strong Ca~II emission being present \\citep{and96}. The inferred electron density in the line formation region is high \\citep{hs88}, and it is likely that the Ca~II 3$^{2}$D level is collisionally de-excited before the forbidden transition can occur. This also appears likely in the case of \\object{W243}, where relatively strong [Ca~II] emission would be required to be detectable amidst the telluric lines.  \\subsection{Conclusions and future work}  Our observations of \\object{W243} show that it has remained in a quiescent, cool-phase state since 2003, with no indication of either further movement redwards or a return to its former hot state. The $\\sim$A3Ia$^{+}$ absorption-line spectrum displays photospheric pulsations and episodic mass loss that appears similar to warm YHGs (e.g. \\object{$\\rho$~Cas}, \\citealt{lobel98} or \\object{W265}, \\citealt{ritchie09}). The high Nitrogen abundance and depletion of Oxygen and Carbon imply that CNO-processed material is on the surface of \\object{W243}, indicating an advanced evolutionary state and significant previous mass-loss. Superimposed on the YHG absorption-line spectrum are emission lines of H, He~I, Ca~II, Fe~II and O~I formed at large radii when material lost from the LBV is ionized by an unseen hot companion. We are unable to strongly constrain the emission-line formation region and the origin of the peculiar, double-peaked Balmer-series profiles or the highly variable He~I emission lines. However, high-resolution infra-red spectroscopy with adaptive optics offers the opportunity to directly probe the line-forming region and thereby determine the recent mass-loss history of \\object{W243}. In addition, abundance studies of the coeval population of both pre- and post-LBV transitional stars in Wd1 will allow the evolutionary state of both \\object{W243} and its evolutionary contemporaries to be determined."
    },
    "0910/0910.2230_arXiv.txt": {
        "abstract": "The universe can be made flat and smooth by undergoing a phase of ultra-slow (ekpyrotic) contraction, a condition achievable with a single, canonical scalar field and conventional general relativity.  It has been argued, though, that generating scale-invariant density perturbations, requires at least two scalar fields and a two-step process that first produces entropy fluctuations and then converts them to curvature perturbations. In this paper, we identify a loophole in the argument and introduce an ekpyrotic model based on a single, canonical scalar field that generates nearly scale-invariant curvature fluctuations through a purely ``adiabatic mechanism\" in which the background evolution is a dynamical attractor. The resulting spectrum can be slightly red with distinctive non-gaussian fluctuations. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.0235_arXiv.txt": {
        "abstract": "We introduce a cosmological model based on the normal branch of DGP braneworld gravity with a smooth dark energy component on the brane. The expansion history in this model is identical to $\\Lambda$CDM, thus evading all geometric constraints on the DGP cross-over scale $r_c$. This well-defined model can serve as a first approximation to more general braneworld models whose cosmological solutions have not been obtained yet. We study the formation of large scale structure in this model in the linear and non-linear regime using N-body simulations for different values of $r_c$. The simulations use the code presented in \\cite{DGPMpaper} and solve the full non-linear equation for the brane-bending mode in conjunction with the usual gravitational dynamics. The brane-bending mode is attractive rather than repulsive in the DGP normal branch, hence the sign of the modified gravity effects is reversed compared to those presented in \\cite{DGPMpaper}.  We compare the simulation results with those of ordinary $\\Lambda$CDM simulations run using the same code and initial conditions. We find that the matter power spectrum in this model shows a characteristic enhancement peaking at $k\\sim0.7\\iMpch$. We also find that the abundance of massive halos is significantly enhanced. Other results presented here include the density profiles of dark matter halos, and signatures of the brane-bending mode self-interactions (Vainshtein mechanism) in the simulations. Independently of the expansion history, these results can be used to place constraints on the DGP model and future generalizations through their effects on the growth of cosmological structure. ",
        "introduction": "\\label{sec:intro}  Braneworld scenarios with infinite extra dimensions have received much interest in recent years, as a possible explanation for the observed accelerated expansion of the universe \\cite{Deffayet01,DeffayetEtal02b}, as well as a way to tackle the cosmological constant problem \\cite{deGrav,deRham}. The simplest such model is the Dvali-Gabadadze-Porrati (DGP) model \\cite{DGP}. In this model, matter lives on a four-dimensional brane embedded in five-dimensional Minkowski space. The gravity action consists of a five-dimensional Einstein-Hilbert term plus an ordinary, four-dimensional term localized on the brane. Gravity thus is five-dimensional on large scales, and reduces to four-dimensional General Relativity on small scales. The transition is given by the \\textit{cross-over scale} $r_c$, which is defined by the ratio of five- and four-dimensional gravitational constants: $r_c = G^{(5)}/2G^{(4)}$.  There are two branches of homogeneous and isotropic solutions in this model, corresponding to the two possible ways of embedding the brane in (asymptotic) five-dimensional Minkowski space. The self-accelerated branch ({\\it sDGP}) has received attention since it leads to a late-time accelerated expansion of the universe without any dark energy or cosmological constant, if $r_c$ is of order the present Hubble horizon $c/H_0$. However, this branch of the model is plagued by a ghost instability when perturbed around the de Sitter solution \\cite{LutyEtal,NicRat,GregoryEtal}. In addition, the expansion history predicted by the self-accelerating DGP model \\cite{Deffayet01} does not appear to fit observations (e.g., \\cite{FangEtal}). The other, so-called normal branch of the theory ({\\it nDGP}), does not have a ghost instability. However, it does not lead to self-acceleration, so that it is necessary to add a cosmological constant (tension) on the brane \\cite{SahniShtanov,LueStarkman,GianantonioEtal,LombriserEtal}, or, more generally a form of stress energy with negative pressure.  Hence, we are interested in generalizing the normal-branch DGP model in a way that allows it to pass expansion history constraints. Much work is ongoing to extend the DGP braneworld scenario \\cite{deRham,galileon,Gabadadze09}, and the proposed models are generally expected to exhibit an expansion history close to $\\Lambda$CDM, while the phenomenology of the modification to GR is expected to be similar to the normal branch DGP model. However, full cosmological solutions have not been obtained yet. Braneworld-inspired {\\it parametrized} generalizations of DGP have been proposed in \\cite{AfshordiEtal,KW}. On the one hand, this approach might well capture features of a yet-unknown full generalized braneworld model. On the other hand, these parametrizations do not constitute a well-defined theory.  Here, we instead opt for considering a well-defined, albeit somewhat contrived model based on DGP, which nevertheless exhibits some of the features of a generalized braneworld model (see \\refsec{braneworld} for a discussion). Our model consists of normal-branch DGP gravity, with a general, smooth quintessence-type dark energy component on the brane tuned so that the resulting expansion history is identical to a $\\Lambda$CDM model. This model which we call \\textit{nDGP+DE} thus satisfies all geometric constraints from, e.g. the cosmic microwave background (CMB) and Supernovae, and leaves the cross-over scale $r_c$ as an almost free parameter of the model. Furthermore, the $r_c \\rightarrow \\infty$ limit of the theory is precisely General Relativity with a cosmological constant. Our main aim in this paper is to study the growth of structure in nDGP+DE in the linear and non-linear regime.  The evolution of linear perturbations in DGP has been studied in \\cite{KoyamaMaartens,SongEtalDGP,CardosoEtal,ScI}. On scales smaller than the horizon and the cross-over scale, the model can be described as an effective scalar-tensor theory. The massless scalar degree of freedom corresponds to displacements of the brane and is called the \\textit{brane-bending mode}. In linear perturbation theory, the effect of the brane-bending mode is simply to rescale the effective gravitational constant for the dynamical potential (the time-time piece of the metric), while the geodesics of photons (i.e. gravitational lensing) are not affected. In the self-accelerating branch of DGP, the brane-bending mode is repulsive, weakening the effective gravitational force, while it is attractive in the normal branch.  As all viable modified gravity models, DGP contains a non-linear mechanism to restore General Relativity in high-density environments. This is achieved by a strong self-coupling of the brane-bending mode, which becomes effective within a characteristic scale, the so-called Vainshtein radius \\cite{Vainshtein72,DeffayetVainshtein02,KoyamaSilva}. The Vainshtein radius depends on the mass considered as well as its configuration. While the prefactor of the self-coupling is model-dependent, the form of the coupling is generic and is expected to be universal to braneworld models \\cite{galileon,ScI}. For typical structures in the Universe, the Vainshtein radius is of cosmological scale. Hence it is necessary to follow the brane-bending mode and its self-interactions when studying the formation of structure in the Universe. In \\cite{DGPMpaper}, we presented simulations of the self-accelerating DGP model which self-consistently solve for the brane-bending mode (see also \\cite{ScII}). Here, we present simulations of two nDGP+DE models, extending the studies of non-linear structure formation presented in \\cite{DGPMpaper} to the normal branch of DGP, where the effects of modified gravity have the opposite sign.  N-body simulations of normal-branch braneworld models have previously been presented in \\cite{KW}. The key difference to our simulations is that while we solve the full brane-bending mode equation, \\cite{KW} employed an approximation where the modified forces are parametrized by an effective density-dependent gravitational constant $G_{\\rm eff}(\\d\\rho)$. Unfortunately, as already pointed out in \\cite{KW}, this approximation does not recover the correct large-distance behavior for the brane-bending mode, so that in principle artefacts of the approximation might appear even on large scales. Furthermore, the approximation will become worse the higher the resolution of the simulation is. We present a quantitative comparison of our simulations with the $G_{\\rm eff}(\\d)$ approximation in \\refapp{KW}.  In addition to the matter power spectrum and halo mass function, we show results on the density profiles of halos, and signatures of the Vainshtein mechanism in the simulation results. The results presented in this paper together with \\cite{DGPMpaper} can serve as a basis for modeling non-linear structure formation in DGP. A model of some of the results presented here and in \\cite{DGPMpaper} in the context of spherical collapse and the halo model will be the subject of a forthcoming paper. Such a model can then be used to constrain DGP and more generalized braneworld models using the wealth of large scale structure observations available (see \\cite{fRcluster} for a first attempt in the case of $f(R)$ gravity).  The paper is structured as follows. In \\refsec{DGP}, we describe the nDGP+DE models in detail, including the evolution of linear perturbations and the Vainshtein mechanism. We describe the simulations in \\refsec{sim}, and present the results in \\refsec{res}. We conclude in \\refsec{concl}. ",
        "conclusions": "\\label{sec:concl}  We have introduced the nDGP+DE model based on the normal branch of DGP with a smooth dark energy component on the brane, which results in an expansion history that is precisely $\\Lambda$CDM. This model can serve as an approximation, especially on sub-horizon scales, to more general braneworld models whose cosmological solutions have not been obtained yet. In addition, the $r_c\\rightarrow \\infty$ limit precisely corresponds to General Relativity plus cosmological constant.  Geometric measurements such as the acoustic scale constrained from the CMB, BAO, and Supernovae observations do not constrain the crossover scale $r_c$ between 4D and 5D gravity in this scenario. The growth of large-scale structure however does offer a sensitive probe of braneworld gravity via the effects of the brane-bending mode which mediates an additional attractive force. In the case of the self-accelerating branch, the brane-bending mode is repulsive, so that the modified gravity effects are sign-flipped.  In order to study the formation of structure in these models, it is necessary to solve the non-linear equations of the brane-bending mode in conjunction with the ordinary gravitational dynamics in N-body simulations \\cite{DGPMpaper}. These non-linear interactions, which are responsible for the Vainshtein mechanism \\cite{Vainshtein72,DeffayetVainshtein02} restoring General Relativity in high density environments, also leave characteristic signatures in the bispectrum in the simulations \\cite{ScI}, which we have found in both sDGP and nDGP simulations. In addition, the non-linearities manifest themselves in a suppression of the brane-bending mode around massive halos.  The matter power spectrum in nDGP+DE shows a characteristic enhancement compared to that of $\\Lambda$CDM, increasing towards smaller scales up to $k\\sim 0.7\\iMpch$, and decreasing on even smaller scales. This is evidenced in a complementary way by the profiles of dark matter halos and their environments (halo-mass correlation function). Similar to what was found in $f(R)$ simulations \\cite{HPMhalopaper}, the inner regions of halos are not affected by the enhanced forces since they assembled at early times. The strongest effects are seen in the near environment of halos, around $2-3\\:r_{200}$, corresponding to the scale where the power spectrum enhancement peaks.  The abundance of massive halos is known to be a sensitive probe of the growth of structure \\cite{WhiteEtal93,EkeEtal98,BorganiEtal01,HPMhalopaper}. We found that indeed massive clusters are $2-4$ times more numerous for the nDGP+DE model with $r_c=500-3000\\:$Mpc. Current cluster abundance measurements (e.g. \\cite{Vikhlinin,RozoEtal09,MantzEtal09}) should be able to constrain the crossover scale $r_c$ to be at least of order Gpc - a constraint which can be placed independently of the specific DGP expansion history.  As the next step in understanding the effects of braneworld gravity on large scale structure, an upcoming paper will detail a model of the matter power spectrum and halo mass function based on spherical collapse and the halo model. Such a model can also serve as an efficient way of dealing with cosmological parameter dependencies when comparing with actual observations \\cite{fRcluster}. In the future, by means of these simulations and a physical model of their results, we will hopefully be able to probe the next generation of braneworld scenarios via their effect on the growth of cosmic structure."
    },
    "0910/0910.2971_arXiv.txt": {
        "abstract": "% Photometric surveys of M31's halo vividly illustrate the wreckage caused by hierarchical galaxy formation.  Several of M31's satellites are being disrupted by M31's tidal field, among them M33 and And I, while other tidal structures are the corpses of satellites already destroyed.  The extent to which M31's satellites have left battle scars upon it is unknown; to answer this we need accurate orbits and masses of the perturbers.  I focus here on M31's 150-kpc-long Giant Southern Stream (GSS) as an example of how these can be determined even in the absence of a visible progenitor.  Comparing N-body models to photometric and spectroscopic data, I find this stream resulted from the disruption of a large satellite galaxy by a close passage about 750 Myr ago.  The GSS is connected to several other debris structures in M31's halo.  Bayesian sampling of the simulations estimates the progenitor's initial mass as $M_\\star = 10^{9.5 \\pm 0.2} M_{\\sun}$, showing it was one of the most massive Local Group galaxies until quite recently. The stream model constrains M31's halo mass to be $(1.8 \\pm 0.5) \\times 10^{12} M_{\\sun}$.  While these small uncertainties neglect several important degrees of freedom, they are likely to remain good even with a more complete model.  Future work on M31's satellites and streams will provide independent constraints on M31's mass, and reveal the shared history of M31 and its halo components. ",
        "introduction": "Despite the term ``Local Group'', M31 has never interacted with the Milky Way and the two galaxy virial radii are far from touching.  So M31 is really just an isolated spiral galaxy that happens to be close enough to study in great detail.  Of course even ``isolated'' galaxies have their satellites; M31 has two nearby ones at around 1\\% of its own stellar mass (M32 and NGC 205), while the spiral M33 has around 10\\% of M31's mass but at such large distance that it is unclear whether it is truly a bound satellite.  M31's disk is far from regular: it is significantly warped, and star formation indicators show a fairly well-defined 10 kpc ring that dominates the total star formation rate. Various authors have attributed features within M31 to the impact of satellites including M32 and NGC 205 \\citep{gordon06}. In fact, \\citet{block06} assigned the entire 10 kpc ring to a Cartwheel-like expanding wave, requiring a collision 250 Myr ago which they attribute to M32.  However, none of the satellite orbits are currently known, so all of these scenarios are highly speculative.  \\begin{figure}[!ht] \\plotone{pandasrev.eps}   % \\caption{ \\label{fig:maps} A $20 \\times 14$ degree view of M31 ($272 \\times 191$ kpc assuming a distance of 780 kpc) from the PAndAS survey \\citep[][maps by A. McConnachie and N. Martin]{mcconnachie09}. The inner map shows RGB branch color while the outer map shows false-color density. An inset shows M33 on the same scale. The image is made from matched-filter maps on a tangent plane projection which are sensitive to RGB stars at the distance of M31. Each individual field is 1 degree across. Visible rectangular artifacts are consequences of chip gaps and variations in field quality. The GSS and the NE and W shelves are all metal-rich features as indicated by the red color of their red giant branch.  Metal-poor features include several dSph galaxies, the tangential arc to the E connecting to the dSph And I (dot in the middle of the GSS), and a radial stream along the NW edge of the survey at the upper right. } \\end{figure}  Photometric surveys of M31's halo have discovered a great deal of stellar debris, some likely from M31's disk and some tracing past or ongoing mergers.  The recent PAndAS survey \\citep[][Figure 1]{mcconnachie09} extends to M33 and will eventually fill in a 150 kpc circle around M31, down to stellar magnitudes of $i_{AB} = 24.5$ at $S/N = 10$. Visible in Figure 1 are the Giant Southern Stream \\citep{ibata01}, as well as newly discovered tangential and radial streams to the E and NW respectively. Ledges or ``shelves'' on the NE and W sides are also apparent. These faint features, only made apparent by counting individual stars, are beginning to illuminate the orbits of objects in M31's halo. ",
        "conclusions": ""
    },
    "0910/0910.5556_arXiv.txt": {
        "abstract": "{} {We present the first magnetic Doppler images of a rapidly oscillating Ap (roAp) star.} {We deduce information about magnetic field geometry and abundance distributions of a number of chemical elements on the surface of the hitherto best studied roAp star, \\hd, using the magnetic Doppler imaging (MDI) code, \\inv, which allows us to reconstruct simultaneously and consistently the magnetic field geometry and elemental abundance distributions  on a stellar surface. For this purpose we analyse time series spectra obtained in Stokes I and V parameters with the SOFIN polarimeter at the Nordic Optical Telescope and recover surface abundance structures of sixteen different chemical elements, respectively ions, including Mg, Ca, Sc, Ti, Cr, Fe, Co, Ni, Y, La, Ce, Pr, Nd, Gd, Tb, and Dy. For the rare earth elements (REE) Pr and Nd separate maps were obtained using lines of the first and the second ionization stage.} {We find and confirm a clear dipolar structure of the surface magnetic field and an unexpected correlation of elemental abundance with respect to this field: one group of elements accumulates solely where the positive magnetic pole is visible, whereas the other group avoids this region and is enhanced where the magnetic equatorial region dominates the visible stellar surface. We also observe relative shifts of abundance enhancement- or depletion regions between the various elements exhibiting otherwise similar behaviour. } {} ",
        "introduction": "\\label{introduction} Many physical processes important for stellar evolution are crucially influenced by stellar magnetic fields. Theoretical and observational frameworks tell us for instance that they influence microscopic diffusion of chemical elements within stellar atmospheres, and they may also redistribute angular momentum or give rise to enhanced hydrodynamical instabilities and thus modify mixing properties in stellar interiors.  Ap stars, representing about 5--10\\% of the upper main sequence stars, exhibit magnetic fields that appear to be highly ordered, very stable and often very strong. Many Ap stars also show strong spectral line profile variations synchronized to stellar rotation. This is attributed to oblique magnetic and rotation axes and to the presence of a non-uniform distribution of chemical elements on their surface within the framework of the oblique rotator model (introduced by Stibbs, \\cite{stibbs1950}). With a few exceptions such inhomogeneities exist only in the atmospheres of A stars with magnetic fields, demonstrating that these fields play a crucial role in their formation and evolution.  The spectra of Ap stars also exhibit a remarkable variety of sometimes unidentified spectral line features. Ryabchikova~et~al. (\\cite{Ryabchik04}) for example find overabundances of some Fe peak elements, which may reach a few dex for Cr and even higher for the REE, especially for their second ionization stage, while other chemical elements are found to be underabundant compared to the solar value. An important subgroup of the Ap stars, the rapidly oscillating Ap stars, in addition exhibits high-overtone, low-degree, non-radial $p$-mode pulsations with periods of 6--21 minutes (Kurtz \\& Martinez \\cite{KM00}). Their observed pulsational amplitudes are modulated with the magnetic field variations, exhibiting maximum amplitude at the phases of highest magnetic field strength (Kurtz \\cite{K82}). This indicates a close connection between the magnetic field and the excitation mechanism of roAp pulsations.  The chemical peculiarities mentioned are attributed to the selective diffusion of ions under the competitive action of radiative acceleration and gravitational settling (Michaud \\cite{michaud1970}) possibly in combination with a weak, magnetically driven wind (see, e.g., Babel \\cite{babel92}).  We focus our research on roAp stars since these objects exhibit a number of new, largely unexplored phenomena related to their pulsational variability while sharing all the properties of other Ap stars. The roAp stars are the only main-sequence stars in which high-overtone p-mode pulsations are easily detected and pulsation modes can be identified. Furthermore, roAp stars represent the only group of pulsating stars for which magnetic and pulsation characteristics can well be constricted observationally. Hence, peculiar pulsating stars are key objects in the investigation of the atmospheric and internal structure of the middle main sequence stars.  Very little is known as yet about the origin and structure of magnetic fields and their connection and interaction with surface abundance patches and pulsation. With the development of high-resolution spectropolarimeters it has become possible to extract the full amount of information about magnetic field and abundance distributions using Stokes parameter observations and applying magnetic Doppler imaging.   We have developed and tested a code for MDI (Piskunov\\,\\&\\,Kochukhov \\cite{PK02}; Kochukhov \\& Piskunov \\cite{KP02}) that allows us to reconstruct simultaneously and consistently the magnetic field vector and abundance distribution of various elements on a stellar surface by taking into account all relevant physics of polarized line formation.  In deriving the surface abundances of numerous light-, iron-peak- and rare earth elements by applying MDI, we get more insight into the relation between the magnetic fields, vertical and horizontal abundance characteristics and pulsation, and thus into the atmospheric structure of Ap and roAp stars. At present surface abundance structures have been derived for only one other roAp star, HD\\,83368 (Kochukhov et al. \\cite{KDP04}).  Our paper is composed in the following way: Sect.\\,\\ref{hd24712} is devoted to the description of the target star, \\hd, on which we comment in Sect.\\,\\ref{observations} regarding our spectropolarimetric observations and the data reduction. Details about MDI, the magnetic field geometry of \\hd, and the surface abundance structures of individual elements can be found in Sect.\\,\\ref{magdi} while Sect.\\,\\ref{discuss} is devoted to the discussion of our results. ",
        "conclusions": "\\label{discuss}  Observing line profiles in several polarization states of stellar radiation and modelling their shapes with magnetic Doppler imaging, we could derive the magnetic field geometry and abundance distributions of numerous elements inhomogeneously distributed over the surface of the roAp star \\hd. This is the first analysis of this kind for a magnetic pulsating star. Our investigation shows a complex picture of the interplay between magnetic field, pulsation and atmospheric inhomogeneities in this prototypical roAp star. However, the low \\vs\\ of the star and the limited spectral resolution of our spectropolarimetric observations did not allow us to reliably recover all but the largest spatial scales of the abundance and magnetic structures. For this reason we will discuss only the gross properties of the chemical spots and the magnetic field, without touching upon the higher-order moments of the surface structures.  We obtain a dominant dipolar magnetic field structure (which might be more complex, though the general picture is certainly dipolar) and elemental abundance patterns that are correlated to this geometry in an unexpected way: instead of abundance enhancement regions on \\emph{both} magnetic poles or around the magnetic equator, we observe huge enhancement or depletion regions around \\emph{either} the phase where the positive magnetic pole is visible \\emph{or} where the magnetic equatorial region dominates the visible stellar surface.  In addition, even more exciting and novel is the fact that these enhancement or depletion regions for various chemical elements are shifted in longitude relative to each other. The center of Mg underabundance as well as that of Sc, Fe, and Ni can be found at phase 0.0, while that of Ti occurs around phase 0.082.  The center of maximum abundance of Tb appears very close to phase $\\varphi$\\,=\\,0.0, that of Co and Gd slightly after $\\varphi$\\,=\\,0.0, around 0.02, Nd spots appear at $\\varphi$\\,=\\,0.04, Pr and Dy at around $\\varphi$\\,=\\,0.06. Y and La are already shifted by 30\\degr\\ or $\\approx$ 0.082 in the direction of increasing phase value. Ce, whose center of maximum abundance lags behind all other elements, is shifted by as much as $\\approx$\\,0.12 or $\\approx$\\,45\\degr.  To investigate the reliability of these longitudinal offsets of the abundance structures we have performed a forward calculation for Y, La, and Ce showing most of the conspicuous shifts. Using the final abundance maps reconstructed for these elements we have synthesized their line profiles, rotating surface structures in the longitudinal direction by small amounts and compared the resulting modification of Stokes\\,$I$ with the typical observational uncertainty. This comparison is illustrated in Fig.~\\ref{Y_La_Ce_rot} (online material). As is evident from this numerical experiment, the high $S/N$ of our observations enables us to easily detect shifts of $\\ga$\\,0.05 for Y and Ce and even smaller shifts for La. Thus the relative longitudinal positions of spots of different elements are precise to within 0.02--0.05 of the rotation period.  \\onlfig{9}{ \\begin{figure*}[htbp] \\begin{center} \\vspace{8mm} \\includegraphics[width=150mm]{./prf_dif.eps} \\caption{Effect of a small rotation of the surface structure of Y, La, and Ce on the Stokes $I$ profiles of the lines of these elements used for mapping. Panels show the difference between the original synthetic profiles and spectra computed with the surface spots shifted by 0.05 (blue) and 0.10 (red) of the rotation period (18 and 36\\degr\\ on the stellar surface, respectively). The grey bar indicates the uncertainty ($\\pm\\sigma$) of the spectra corresponding to the $S/N$ of individual rotation phases. The difference spectra are offset vertically for display purposes. The horizontal scale is arbitrary. } \\label{Y_La_Ce_rot} \\end{center} \\end{figure*} }  The diverse surface abundance behaviour observed in \\hd\\ most likely gives us information on how the magnetic field influences the chemical diffusion of the various species within the stellar atmosphere. Accordingly we tried to correlate our results with simple atomic parameters (e.g., atomic weight or excitation potential), but did not find any clear evidence for such a connection. This gives rise to the speculation, as the effect is clearly present, that there is a more complex correlation of the various atomic characteristics influencing chemical diffusion through stellar atmospheres in presence of a magnetic field. Thus these relative shifts in longitude of the mapped elements provide valuable constraints for diffusion models that try to explain how the magnetic field and other hydrodynamical processes regulate atmospheric diffusion.   Moreover, we obtain shifts in latitude of abundance enhancement regions exhibiting otherwise comparable accumulation tendencies. Sc, Ti, Fe, and Ni in this context appear in similar latitudes, whereas Mg and Co seem to be shifted to lower latitudes by 30\\degr\\ and 45\\degr, respectively, compared to the above mentioned elements. Among the elements belonging to the group that is enhanced around the positive magnetic pole, Tb and Y seem to cover the highest latitudes, Ce and Pr are found to be shifted to lower latitudes by around 15\\degr, and La, Nd, Gd, and Dy spots are shifted even more southwards by $\\approx$ 30\\degr.  This finding helps to explain the phase shift between pulsation curves of spectral lines belonging to different REE elements, for instance, Pr and Nd, that we found in the recent study by Ryabchikova~et~al. (\\cite{Ryabchik07}). According to the NLTE analysis of Nd\\,\\ii,\\iii\\ and Pr\\,\\ii,\\iii\\ line formation in the atmosphere of HD\\,24712 (Mashonkina~et~al. \\cite{nd-NLTE,pr-NLTE}), the lines of both elements are formed at nearly the same atmospheric layers. However, the pulsation analysis of Ryabchikova~et~al. (\\cite{Ryabchik07}) shows that they are shifted by 0.1 in the pulsation phase relative to each other.  According to the theoretical calculations by Sousa\\,\\&\\,Cunha (\\cite{SC08}) and numerical simulations by Khomenko\\,\\&\\,Kochukhov (\\cite{KK09}), the amplitude and phase distributions with height in roAp atmospheres should depend on the local magnetic field strength and inclination. Slightly different location of Pr and Nd spots relative to the magnetic field pole may therefore be a reason for the pulsation amplitude and phase difference observed in Pr and Nd lines. This illustrates the potential and importance of MDI when interpreting pulsational observations of roAp stars.   Vertically stratified abundances are very likely present in the atmosphere of \\hd\\ (e.g., Ryabchikova et al. \\cite{RKB08}). In particular, a major discrepancy between two ionization states of the same REE element is attributed to stratification (Mashonkina et al. \\cite{pr-NLTE}). To estimate the impact of stratification, we performed an analysis of the vertical gradient of Fe abundance (L\\\"uftinger et al. \\cite{Lueftinger04}) during the phase where the magnetic equatorial region dominates the stellar hemisphere and also where the positive magnetic pole is visible, choosing Fe lines that are most sensitive to stratification. Fe turned out to be significantly stratified, with abundance changing by $\\approx$\\,2~dex between the upper and lower parts of the atmosphere of \\hd. We also found a slight difference in the location of the abundance jump comparing results from the two extreme phases. In order to keep MDI results influenced as little as possible by the stratification effects, we chose lines for mapping that are hardly affected by a changing vertical distribution. It is still possible however that part of the horizontal abundance structure we find is due to variation of the chemical stratification profile across the stellar surface.   Currently, there are no theoretical diffusion models that can provide quantitative explanation of the abundance maps derived in our Doppler imaging study of \\hd. For this reason we are not able to compare the results  of our investigation to extensive theoretical predictions for individual elements in the atmospheres of Ap stars. The standard diffusion calculations consider chemical transport processes in the presence of axisymmetric magnetic field and thus anticipate abundance inhomogeneities to be axially symmetric themselves (Michaud et al. \\cite{MCM81}). These predictions of the simple diffusion theory are inconsistent with the MDI results obtained for \\hd. We find a nearly dipolar magnetic field geometry coexisting with diverse abundance distributions, frequently shifted relative to the magnetic field axis. A similar inconsistency between a simple magnetic field topology and complex abundance distributions was found in other recent DI studies of Ap stars (Kochukhov et al. \\cite{KPI02,KDP04}). It is possible that diffusion processes are sensitive to the small scale magnetic structures that could be resolved only with the full Stokes vector MDI, such as the study carried out by Kochukhov et al. (\\cite{KBW04}) for 53~Cam. On the other hand one can also suspect that the current simple diffusion theory lacks certain physical processes other than the magnetic field, that can influence or even drive chemical spot formation. Recent discovery of the evolving chemical clouds in the atmosphere of \\textit{non-magnetic} chemically peculiar stars (Kochukhov et al. \\cite{KAG07}) has demonstrated that the surface structure formation in the upper main sequence stars is not necessarily driven by the magnetic field alone.   Anticipating further development of theoretical aspects, we provide new information for the diffusion properties within the atmosphere of a cool Ap star, obtain new insights in the atmospheric structure and the field geometry, the origin and interplay of abundance variations, pulsation and magnetic fields in these unique stellar laboratories."
    },
    "0910/0910.2695_arXiv.txt": {
        "abstract": "I consider couplings between the dark energy and dark matter sectors. I describe how the existence of an adiabatic regime, in which the dark energy field instantaneously tracks the minimum of its effective potential, opens the door for a catastrophic instability. This {\\it adiabatic instability} tightly constrains a wide class of interacting dark sector models. This talk was presented at, and will appear in the proceedings of the DPF-2009 conference. ",
        "introduction": "Modern cosmology demands the existence of two new components to the energy budget of the universe. Dark matter is necessary for structure formation to occur properly, and to account for numerous observations, such as gravitational lensing, the comic microwave background (CMB) and galaxy rotation curves. The discovery of cosmic acceleration requires a second source of new physics, which may come in the form of a modification to general relativity, but which is perfectly consistent with a cosmological constant or a dynamical dark energy component~(for reviews see~\\cite{Copeland:2006wr,Linder:2008pp,Frieman:2008sn,Silvestri:2009hh,Caldwell:2009ix}).  At a phenomenological level, this description is a remarkable fit to all current observations. However, at the level of fundamental physics the existence of dark matter and dark energy, comprising the vast majority of the contents of the universe, poses a critical challenge: how do these components fit into our microphysical theories of matter and energy? One way to search for such relationships is to explore observable consequences of interactions between dark components and baryonic matter. Examples of this are the search for the annihilation products of dark matter and collider searches for dark matter.  Alternatively, if dark matter and dark energy are to fit into a unified description, we might expect~\\cite{combined} that there would also be interactions between them. In this talk I briefly explored some consequences of such interactions, focusing on a possible catastrophic instability - the {\\it adiabatic instability} - and a number of constraints on stable models. As with all proceedings with page limits, thorough referencing is impossible, and I have therefore referenced mostly review articles, and the papers I actually referred to in the talk. ",
        "conclusions": "I have briefly discussed a broad class of models in which dark matter is coupled to a dark energy component, assumed to be responsible for the acceleration of the universe. Within these models, I have discussed a possible instability - the adiabatic instability - that arises in a range of cosmological and astrophysical settings, and which rules out a set of parameter values. I have discussed specific examples of models in which this instability is active, and it is worth noting that one may carry out similar analyses to constrain subclasses of chameleon and MaVan models."
    },
    "0910/0910.1375_arXiv.txt": {
        "abstract": "Recently we examined a large number of points in a 19-dimensional parameter subspace of the CP-conserving MSSM with Minimal Flavor Violation.  We determined whether each of these points satisfied existing theoretical, experimental, and observational constraints.  Here we discuss the properties of the parameter space points allowed by existing data that are relevant for dark matter searches. ",
        "introduction": "The MSSM has a large parameter space; this raises the question of how well we know its properties aside from specific SUSY breaking scenarios such as mSUGRA, AMSB, GMSB, etc. In an attempt to address this question, we performed scans of a 19-dimensional subspace of the full 100+ parameter (R-parity conserving) MSSM (sometimes referred to as the ``phenomenological MSSM'') and applied a comprehensive set of theoretical, experimental, and observational constraints, thereby obtaining a large set of ``models'' (parameter space points) which are consistent with existing data\\cite{us}. The particulars of and results from this scan are discussed in other presentations at this conference\\cite{tom,john}.  Here we will be concerned with the implications of this scan for dark matter, addressing the properties of LSPs in the set of allowed MSSM models generated, and in particular their signatures in direct detection and indirect detection experiments. ",
        "conclusions": ""
    },
    "0910/0910.3693_arXiv.txt": {
        "abstract": "Cosmic microwave background studies of non-Gaussianity involving higher-order multispectra can distinguish between early universe theories that predict nearly identical power spectra. However, the recovery of higher-order multispectra is difficult from realistic data due to their complex response to inhomogeneous noise and partial sky coverage, which are often difficult to model analytically. A traditional alternative is to use one-point cumulants of various orders, which collapse the information present in a multispectrum to one number.  The disadvantage of such a radical compression of the data is a loss of information as to the source of the statistical behaviour. A recent study by Munshi \\& Heavens (2009) has shown how to define the skew spectrum (the power spectra of a certain cubic field, related to the bispectrum) in an optimal way and how to estimate it from realistic data. The skew spectrum retains some of the information from the full configuration-dependence of the bispectrum, and can contain all the information on non-Gaussianity. In the present study, we extend the results of the skew spectrum to the case of two degenerate power-spectra related to the trispectrum. We also explore the relationship of these power-spectra and cumulant correlators previously used to study non-Gaussianity in projected galaxy surveys or weak lensing surveys. We construct nearly optimal estimators for quick tests  and generalise them to estimators which can handle realistic data with all their complexity in a completely optimal manner. Possible generalisations for arbitrary order are also discussed. We show how these higher-order statistics and the related power spectra are related to the Taylor expansion coefficients of the potential in inflation models, and demonstrate how the trispectrum can constrain both the quadratic and cubic terms. ",
        "introduction": "The inflationary paradigm which solves the flatness, horizon and monopole problem makes clear predictions about the generation and nature of density perturbations \\citep{Guth81,Staro79,Linde82,AlSt82,Sato81}. Inflationary models predict the statistical nature of these fluctuations, which are being tested against data from a range of recent cosmological observations, including the recently-launched all-sky cosmic microwave background (CMB) survey Planck\\footnote {http://www.rssd.esa.int/index.php?project=Planck}. Various ground-based and space-based observations have already confirmed the generic predictions of inflation, including a flat or nearly flat universe with nearly scale-invariant adiabatic perturbations at large angular scales. Several planned missions are also targeting the detection of the gravitational wave background through polarisation experiments - another generic prediction of inflationary models. The other major prediction of the inflationary scenarios is the nearly Gaussian nature of these perturbations. In the standard slow-roll paradigm, the scalar field responsible for inflation fluctuates with a minimal amount of self interaction which ensures that any non-Gaussianity generated during the inflation through self-interaction would be small \\citep{Salopek90,Salopek91,Falk93,Gangui94,Acq03,Mal03}. See \\cite{Bartolo06} for a recent review and more detailed discussion. Any detection of non-Gaussianity would therefore be a measure of self-interaction or non-linearities involved, which can come from various alternative scenarios such as curvaton mechanism, warm inflation, ghost inflation as well as string theory inspired D-cceleration and Dirac Born Infield (DBI) models of inflation \\citep{lindemukha,GuBeHea02,Lyth03}.  Early observational work on detection of primordial non-Gaussianity, from COBE \\citep{Komatsu02} and MAXIMA \\citep{Santos} was followed by much more accurate analysis with WMAP\\footnote{http://map.gsfc.nasa.gov/}\\citep{Komatsu03,Crem07a,Spergel07}. Optimised 3-point estimators were introduced by \\citet{Heav98}, and have been successively developed \\citep{KSW,Crem06,Crem07b,SmZaDo00,SmZa06}. Indeed, now an estimator for $f_{NL}$ which saturates the Cramer-Rao bound has been found, capable of treating partial sky coverage and inhomogeneous noise \\citep{SmSeZa09}. The recent claim of a detection of non-Gaussianity in WMAP data \\citep{YaWa08} has given a tremendous boost to the study of  primordial non-Gaussianity, as it can lift the degeneracy between various early-universe theories which predict nearly the same primordial power spectrum. Most detection strategies focus on lowest order in non-Gaussianity, i.e. the bispectrum or three-point correlation functions \\citep{KSW,Crem03,Crem06,MedeirosContaldi06,Cabella06,Liguori07,SmSeZa09}. This is primarily because of a decrease in signal-to-noise as we move up in the hierarchy of correlation functions - the higher-order correlation functions are more dominated by noise than their lower-order counterparts. Another related complication arises from the necessity to optimise such estimators, and the impact of inhomogeneous noise and partial sky coverage is always difficult to include in such estimates.  The recent study by \\cite{MuHe09} suggested the possibility of finding optimised cumulant correlators associated with higher-order multi-spectra in the context of CMB studies. These correlators are well-studied in the context of projected surveys such as projected galaxy surveys or in the context of weak lensing studies using simulated maps \\citep{Mu00,MuJe00,MuJe01}. Early studies involving cumulant correlators focused mainly on understanding gravity-induced clustering in collisionless media and were widely employed in many studies involving numerical simulations. Cumulant correlators are multipoint correlation functions, collapsed to two points Although they are two-point correlators, they carry information on the corresponding higher-order correlation functions. Due to their {\\em reduced} dimensionality they do not carry all the information that is encoded in higher-order correlation function, but they carry more than their one-point counterparts, namely the moments of the probability distribution function which are often used as clustering statistics \\citep{BernardReview}.  One of the reasons to go beyond the lowest order in non-Gaussianity was pointed out by many authors, including. e.g., \\cite{RiqSper}. At smaller angular scales, the secondary effect may dominate \\citep{SperGold99a,SperGold99b}, in direct contrast to larger angular scales where the anisotropies are generated mainly at the surface of the last scattering. These secondary perturbations are produced by interaction of CMB photons at much lower redshift with the intervening large-scale matter distribution. Such effects will be directly observable with Planck. The deflection of CMB photons by the large-scale mass distribution offers the possibility of studying the statistics of density perturbations in an unbiased way and provide clues to growth of structure formation for most part of the cosmic history. Weak lensing of the CMB can provide valuable information for constraining neutrino mass, dark energy equation of state and also has the potential to assist detection of primordial gravitation waves through CMB polarisation information; see e.g. \\cite{LewChal} for a recent review. However weak lensing studies using the CMB need to address the contamination produced by other secondaries such as the thermal Sunyaev Z\\'eldovich (tSZ) effects and kinetic Sunayev Z\\'eldovich effect (kSZ) as well as by point sources, Although it is believed that these contaminations are not so important in case of polarisation studies. It was pointed out in \\cite{RiqSper} that a real space statistic such as $\\langle \\delta T^3(\\oh) \\delta T(\\oh') \\rangle_c$ (which is a cumulant correlator of order four) can be used to separate the kSZ effect or Ostriker-Vishniac (OV) effect from the lensing effect as the lensing contribution cancels out at the lowest order. It was shown that, in addition to quantifying and controlling the kSZ contamination of lensing statistics such statistics could as well play a very important role in providing new insight into the history of reionization. In a completely different context it was shown that this estimator also has use for testing models of primordial non-gaussianity using redshifted 21cm observations \\citep{Coo21,Cooray8}. While cumulant correlators at third order $\\langle \\delta T^2(\\oh) \\delta T(\\oh') \\rangle_c$ can provide information regarding the non-Gaussianity parameter $f_{\\rm NL}$, fourth-order statistics such as $\\langle \\delta T^2(\\oh) \\delta T^2(\\oh') \\rangle_c$ can go beyond the lowest order non-Gaussianity by putting constraints on the next-order parameter $g_{\\rm NL}$ (to be introduced in later sections), albeit at lower signal-to-noise. However with ongoing CMB missions such as Planck the situation will improve and developing optimal methods for such higher-order cumulant correlators is the first step in this direction. There are now several studies which provide independent estimates of $f_{\\rm NL}$ however we still lack such constraints for $g_{\\rm NL}$. This clearly is related to the fact that in typical models $g_{\\rm NL}$ is expected to be small,  $g_{\\rm NL} \\le r /50 $, where $r$ is the scalar-to-tensor ratio \\citep{Seery07}. We also note that various studies have pointed out the link between the non-Gaussianity analysis and the estimators which test the anisotropy of the primordial universe. We plan to address these issues in a related publication. Several inflationary models provide direct consistency relations between $f_{\\rm NL}$ and $g_{\\rm NL}$, e.g. $g_{NL} = (6f_{\\rm NL}/5)^2 $ (in some publications it is also denoted as $f_2$ or $\\tau_{NL}$ e.g. \\citep{huOka02,Coo21,Seery07}). Testing of these consistency relations can give valuable clues to the mechanism behind the generation of initial perturbations. However, a difficulty for methods designed to detect non-Gaussianity in the CMB is that other processes can contribute, such as gravitational lensing, unsubtracted point sources, and imperfect subtraction of galactic foreground emission  discussed by, e.g., \\citet{GoldbergSpergel99,CoorayHu,VerdeSperg02,Castro04,Babich08}.  While the main motivation in this work is to study the primordial trispectrum, we note that mode-mode coupling resulting from weak lensing of the CMB produces a trispectrum which has been studied using power-spectra associated with $\\langle \\delta T(\\oh)^2 \\delta T(\\oh')^2 \\rangle_c$. Lensing studies involving the CMB can achieve higher signal-to-noise ratio at the level of the bispectrum if we use external data sets to act as a tracers of large-scale structure.  However, the lowest order at which the internal detection of CMB lensing is possible is the trispectrum.   There have been some attempts to detect non-gaussianity at the level of the trispectrum using e.g. COBE 4yr data release \\citep{Kunz01}, BOOMERanG data \\citep{Troia03} and more recently using WMAP 3 year data \\citep{Spergel07}.  The layout of this paper is as follows. The section \\textsection2 we introduce the concept of the higher-order cumulant correlators and how they are linked to corresponding correlation functions in real space. In \\textsection3 we discuss the harmonic transforms of the cumulant correlators and their relations to the corresponding multi-spectra. We also discuss estimators based on pseudo-$C_l$ (PCL) estimators used for power spectrum estimation and generalise them to two-point cumulant correlators at higher order in this section. In \\textsection4 we briefly discuss the ``{\\em local}'' models for initial perturbations and the resulting trispectrum. These models are then used in \\textsection4 to optimise the power spectra associated with the multi-spectra. This approach is nearly optimal, describing the mode-mode coupling using the ``{\\em fraction of sky}'' approach familiar from other studies. In \\textsection5 we present an estimator which is nearly optimal and can take into account partial sky coverage as well as realistic inhomogeneous noise resulting from the scanning strategy. This estimator can work directly with any specific theoretical model for primordial trispectra, based on concepts of matched filtering, and generalises the results obtained previously by \\cite{MuHe09}. We present results both for one-point estimators and also generalise them to two corresponding power spectra. \\textsection6 is devoted to adding relevant linear correction terms to these estimators in the absence of spherical symmetry. Next, in \\textsection7 we introduce inverse covariance weighting and design a trispectrum estimator which is optimal for arbitrary scanning strategy. The estimator used in this section will also be useful also in estimating secondary non-Gaussianity. For homogeneous noise and all-sky coverage the estimators are identical. Finally, \\textsection8 is devoted to concluding discussions and future prospects. ",
        "conclusions": "In the near future the all-sky Planck satellite will complete mapping the CMB sky in unprecedented detail, covering a huge frequency range. The cosmological community will have the opportunity to use the resulting data to constrain available theoretical models. While the power spectrum provides the bulk of the information, going beyond this level will lift degeneracies among various early universe scenarios which otherwise have near identical power spectra. The higher-order spectra are the harmonic transforms of multi-point correlation functions, which contain information which can be difficult to extract using conventional techniques. This is related to their complicated response to inhomogeneous noise and sky coverage. A practical advance is to form collapsed two-point statistics, constructed from higher-order correlations, which can be extracted using conventional power-spectrum estimation methods.  We have specifically studied and developed three different types of estimators which can be employed to analyse these power spectra associated with higher-order statistics. The MASTER-based approach  \\citep{Hiv} is typically employed to estimate pseudo-$C_l$s from the masked sky in the presence of noise; also see \\cite{Efs1,Efs2}. These are unbiased estimators but the associated variances and scatter can be estimated analytically with very few simplifying assumptions. We extend these estimators to study higher-order correlation functions. We develop estimators for $C_l^{(2,1)}$ for the skew-spectrum (3-point) as well as $C_l^{(3,1)}$ and $C_l^{(2,2)}$ which are power spectra of fields related to the trispectrum, or kurt-spectrum (4-point). The removal of the Gaussian contribution is achieved by applying the same estimators to a set of Gaussian simulations with identical power spectra and subtracting.  As a next step we generalised the estimators employed by \\cite{YKW,YaWa08,Yadav08} and others to study the kurt-spectrum. These methods are computationally expensive and can be implemented using a Monte-Carlo pipeline which can generate 3D maps from the cut-sky harmonics using radial integrations of a target theoretical model. The Monte-Carlo generation of 3D maps is the most computationally expensive part and dominates the calculation. The technique nevertheless has been used extensively, as it remains highly parallelisable and is optimal in the presence of homogeneous noise and near all-sky coverage. The corrective terms involves linear and quadratic terms for lack of spherical symmetry due to inhomogeneous noise and partial sky coverage. These terms can be computed using a Monte-Carlo chain. We also showed that the  radial integral involved at the three-point analysis needs to extended to a double-integral for the trispectrum. The speed of this analysis depends on how fast we can generate non-Gaussian maps. It is possible also to use the same formalism to study cross-contamination from other sources such as point sources or other secondaries while also determining primordial non-Gaussianity. The analysis also allows us to compute the overlap or degeneracy among various theoretical models for primordial non-Gaussianity.   Finally we presented the analysis for the case of estimators which are completely optimal even in the presence of inhomogeneous noise and arbitrary sky coverage \\citep{SmZa06}. Extending previous work by \\cite{MuHe09} which concentrated only on the skew-spectrum, we showed how to generalise to the trispectrum. This involves finding a fast method to construct and invert the covariance matrix $C_{lml'm'}$ in multipole space. In most practical circumstances it is possible only to find an approximation to the exact covariance matrix, and to cover this we present analysis for an approximate matrix which can be used as instead of $C^{-1}$. This makes the method marginally suboptimal but it remains unbiased. The four-point correlation function also takes contributions which are purely Gaussian in nature. The subtraction of these contributions is again simplified by the use of Gaussian Monte-Carlo maps with the same power spectrum. A Fisher analysis was presented for the construction of the error covariance matrix, allowing joint estimation of trispectra contributions from various sources, primaries or secondaries. Such a joint estimation give us fundamental limits on how many  sources of non-Gaussianity can be jointly estimated from a specific experimental set up. A more detailed Karhuenen-Lo\\'eve eigenmode analysis will be presented elsewhere.  The detection of non-primordial effects, such as weak lensing of the CMB, the kinetic SZ effect, the Ostriker-Vishniac effect and the thermal SZ effect can provide valuable additional cosmological information \\citep{RiqSper}.    The detection of CMB lensing at the level of the bispectrum needs external data sets which trace large-scale structures, but the trispectrum offers an internal detection, albeit with reduced S/N,  providing a valuable consistency check.  At the level of the bispectrum, primordial non-Gaussianity can for many models be described by a single parameter $f_{NL}$. The two degenerate kurt-spectra (power spectra related to tri-spectrum) we have studied at the 4-point level require typically two parameters such as $f_{NL}$ and $g_{NL}$. Use of the two power spectra will enable us to put separate constraints on $f_{NL}$ and $g_{NL}$ without using information from lower-order analysis of bispectrum, but they can all be used in combination. Clearly at even higher-order more parameters will be needed to describe various parameters ($f_{NL}$, $g_{NL}$, $h_{NL}$, \\dots) which will all be essential in describing degenerate sets of power spectra associated with multispectra at a specific level. Note we should keep in mind that higher-order spectra will be progressively more dominated by noise and may not provide useful information beyond a certain point.  The power of the estimators we have constructed largely depends on finding a techniques to simulate non-Gaussian maps with a specified bispectrum and trispectrum. We have discussed the possibility of using our technique to generate all-sky CMB maps with specified lower-order spectra, i.e. power spectrum, bispectra and trispectra. These will generalise previous results by \\cite{SmZa06} which can guarantee a specific form at bispectrum level.  While we have primarily focused on CMB studies, our estimators can also be useful in other areas e.g. studies involving 21cm studies, near all-sky redshift surveys and weak lensing surveys. The estimators described here can be useful for testing theories for primordial and/or gravity-induced secondary trispectrum using such diverse data sets. In the near future all-sky CMB experiments such as Planck can provide maps covering a huge frequency range and near-all sky coverage. The estimators described here can play a valuable role in analysing such maps."
    },
    "0910/0910.3977_arXiv.txt": {
        "abstract": "We present the results of new CCD photometry for the contact binary BX Peg, made during three successive months beginning on September 2008. As do historical light curves, our observations display an O'Connell effect and the November data by themselves indicate clear evidence for very short-time brightness disturbance. For these variations, model spots are applied separately to the two data set of Group I (Sep.--Oct.) and Group II (Nov.). The former is described by a single cool spot on the secondary photosphere and the latter by a two-spot model with a cool spot on the cool star and a hot one on either star. These are generalized manifestations of the magnetic activity of the binary system. Twenty light-curve timings calculated from Wilson-Devinney code were used for a period study, together with all other minimum epochs. The complex period changes of BX Peg can be sorted into a secular period decrease caused dominantly by angular momentum loss due to magnetic stellar wind braking, a light-travel-time (LTT) effect due to the orbit of a low-mass third companion, and a previously unknown short-term oscillation. This last period modulation could be produced either by a second LTT orbit with a period of about 16 yr due to the existence of a fourth body or by the effect of magnetic activity with a cycle length of about 12 yr. ",
        "introduction": "For BX Peg, Lee at al. (2004a, hereafter Lee04a) remains the most recent comprehensive study. These authors showed that historical light curves of BX Peg, displaying year-to-year light variability, can all be explained by introducing a single dark spot on the more massive secondary star and found that the orbital period has varied due to a periodic oscillation overlaid on a continuous period decrease. They concluded that the periodic $O$--$C$ residuals could not be produced by spot activity (Maceroni \\& van't Veer 1994) and were not locked into the light variation as required by Applegate (1992). Thus, the only phenomenon that could be responsible was a third body with a period of 52.4 yr and a limiting mass of $M_3 \\sin i_3$=0.26 $M_\\odot$ and the hypothetical companion became the default explanation.  Lee04a also suggested that the timing residuals indicated an additional short-term oscillation with a period of about 12 yr and a semi-amplitude of about 0.002 d. Since then, many new photoelectric and CCD timings of minimum light have been reported and should now be sufficiently numerous to test this possibility meaningfully. In this article, we present a detailed study of the $O$--$C$ diagram of BX Peg together with a new light-curve synthesis. ",
        "conclusions": "The negative coefficients of the quadratic terms in equations (2) and (3) yield a continuous period decrease with a rate of $-$9.8 $\\times$10$^{-8}$ d yr$^{-1}$, which can be explained either by conservative mass transfer from the more massive cool star to its less massive hot component or by angular momentum loss (AML) due to a magnetic stellar wind. Under the assumption of conservative mass transfer, the transfer rate is about 7.0$\\times$10$^{-8}$ M$_\\odot$ yr$^{-1}$. This value is 3.5 times greater than a rate of 2.0$\\times$10$^{-8}$ M$_\\odot$ yr$^{-1}$ calculated by assuming that the cool secondary transfers mass to the hot primary component on a thermal time scale. Therefore, the alternative mechanism, AML caused by magnetic braking rooted in the convective zone, seems more likely. From an approximate formula given by Guinan \\& Bradstreet (1988), the period decrease rate is calculated to $-$4.1, $-$5.8, and $-$8.7 (in units of 10$^{-8}$ d yr$^{-1}$) for $k^2$=0.07, 0.10, and 0.15, respectively. The last value might be a good approximation to the gyration constant $k^2$ of BX Peg. It is also possible that the correct explanation of the secular period change is some combination of non-conservative mass transfer and AML but, at present, it seems that AML should be the dominant contributor.  This last cautionary remark actually has some independent support. Consider the pair of W-type systems defined observationally by Binnendijk (1970), V829 Her (Erdem \\& \\\"Ozkarde\\c s 2006) and V781 Tau (Yakut et al. 2005). Their individual masses and radii and the photospheric temperature difference as well as the systemic mass ratio and period are the same values within 1 \\% yet the algebraic signs of the secular period changes are different. This is not a unique case. The same situation exists for V417 Aql (Qian 2003; Lee et al. 2004b) and BB Peg (Kalomeni et al. 2007). If there is little mechanical and radiometric difference between two such binaries which behave differently dynamically, there must be a particular evolutionary mechanism that distinguishes them or else the seeming secular period changes are not really unbounded and are just long-term cyclical effects. The rigorous association of a positive period change with a particular sense of mass transfer and the mirror association of a negative period change with the other sense of mass transfer and with AML appear to be too didactic.  The 12-yr period modulation, shown in the third panel of Figure 4, could operate in at least one component since each has a convective envelope and thus may be magnetically active (Applegate 1992;  Lanza et al. 1998). With the values of $K^{\\prime}$ and $P^{\\prime}$, the parameters of an Applegate model were calculated for both components and appear in Table 7, where $\\Delta m_{\\rm rms}$ denotes a bolometric magnitude difference relative to the mean light level of BX Peg converted to magnitude scale with equation (4) in the paper of Kim et al. (1997). The variations of the gravitational quadrupole moment $\\Delta Q$ correspond to typical values for contact binaries and the required light variation associated with each component is within the value ($\\Delta L/L_{\\rm {p,s}} \\sim 0.1$) proposed by Applegate. This consistency indicates that Applegate's mechanism could possibly function in both component stars. There is no {\\it a priori} reason to require a sine-type behavior for a magntic cycle in a star. If Applegate's mechanism is the main cause of the cyclical variation, his model requires the brightness variation to vary in phase with the period modulation of Figure 4.  In order to check this possibility and to study the long-term light variations of BX Peg, we measured the light levels at four different phases (Max I, Min I, Max II and Min II) for our new light curves and for the archival measurements of Zhai \\& Zhang (1979), Samec (1990) and Lee04a.  The latter are taken from Table 6 of Lee04a and the comparisons are given in Table 8.  All data sets are referred to the same comparison star (${\\rm BD}+25^{\\rm o}4584$). The brightness variations of $\\Delta B$ and $\\Delta V$ are plotted in Figure 6, where the seasonal means were subtracted from the grand mean for all seasons giving the mean seasonal differences in the natural systems. The year and epoch for each data set were calculated by averaging the starting and ending HJDs of the observations. In Figure 6, the third and fourth panels represent the sine term of the $C_2$ ephemeris and the $\\tau_4$ orbit of the $C_3$ ephemeris, respectively, averaged over each observing season and the arrows indicate the mean epoch of each seasonal light curve. There are only 5 epochs for light curves but, even as few as they are, they do not conform to this prediction of the Applegate mechanism.  This binary presents a conflict for observers.  Without the minimum monitoring which has been so fruitful, we would be ignorant of the time-scales and semi-amplitudes of the 50-year cycle and of the shorter one as well.  Perusal of the present Figures 3 through 5 and of Figures 4 and 5 in Lee04a shows that the 50-year cycle has become much better delineated over only 5 years.  The same conclusion cannot be drawn for the shorter cycle.  We have already remarked that neither the sinusoid nor a fourth-companion waveform track the residuals well.  Many of the minimum timings are much less precise than those from LOAO and their noisy appearance is readily evident in all the figures.  Residuals as large as $-$3.7 minutes appear in Table 4, a value that is almost 1 \\% of the Keplerian period and more than 4 \\% of the entire eclipse width.  Such discrepancies probably do not signify high-frequency spot activity; at least there is no independent evidence for such a possibility.  Rather, they must be the result of clerical or observational errors or of greatly-undersampled timing determinations."
    },
    "0910/0910.1972_arXiv.txt": {
        "abstract": "We study equilibrium solutions of a boson star in mean field approximation using a general relativistic framework. Dynamics of the constituent matter in such a star is described by a scalar field. The statistical aspect of the matter equilibrium is investigated by incorporating the effect of finite chemical potential in equation of state. We analyze the two possible situations where  the scalar field is either a composite one -suitable for describing diquark-condensates or a fundamental. We derive a generalized set of Tolman-Oppenheimer-Volkov (TOV) equations to incorporate the metric dependence of chemical potential in general relativity. It is demonstrated that the introduction of finite chemical potential can lead to a new class of the solutions where the maximum mass and radius of a star change in a significant way. In general the density profile has node like structures due to introduction of finite chemical potentials. It is shown that these nodes can be avoided by applying constraints on the values of the central pressure and central chemical potential. This in turn can reduce the parameter space available for the stable solutions. We also discuss the case when the boson star is made of diquark condensates. It is shown that when the self-interaction of the condensate is negligible, the typical energy density of the star is too high to have a boson star of diquark condensates. Our results indicate, even after the effect of self-interaction is incorporated, the stable equilibrium can occur only in a low-density and high-pressure regime. ",
        "introduction": "The boson stars are very fascinating objects as their self-gravity is not balanced by the degeneracy pressure like the other compact stars such as a white dwarf or a neutron star, but the Heisenberg uncertainty play a crucial role in their stability (see Ref. \\cite{sch} for a recent review). Wheeler was first to consider the boson star in 1955 when he studied the stable equilibrium generated by the self-gravitating photons called {\\it geons} \\cite{wheel}. Later Kaup analyzed a new class of boson stars as a self-consistent solutions of Einstein-Klein-Gordon equations in Ref. \\cite{kaup}. Since then the most of the boson star studies focus on the the stars made of self-gravitating scalar-field \\cite{sch, rev2}. This is because there are good reasons to believe that there exists a some fundamental scalar field in nature even though no experiment has detected it so far. The solutions found by Kaup in Ref. \\cite{kaup} were shown to describe the ground state of boson stars without using any perfect-fluid approximation or some approximation to equation of state by Ruffini and Bonazzola \\cite{rb}. However in this work the self-interaction of the scalar-field was ignored. Subsequently it was shown by Colpi et. al. \\cite{shapiro} that inclusion of even a very small interaction term can significantly alter the solutions obtained in Refs. \\cite{kaup,rb}. It turns out that the inclusion of a small interaction term in the scalar field Lagrangian give rise to an additional pressure that would help the Heisenberg uncertainty to counter the self-gravity. This results in a boson star with much larger  mass than ones reported in Refs. \\cite{kaup,rb}. Also the issues of dynamical stability  of boson stars have been rigorously analyzed by applying perturbative methods \\cite{stab}. Also are studies that show that there can be a boson star in the galactic center \\cite{gallaxy-bs}. At present a very intriguing situation persists in the field: on the one hand there is no observational evidence which can indicate that a boson star exists, on the other hand their existence seem to be a reasonable consequence of well-tested physical theories like general relativity and Klein-Gordon equations. This should be a good enough reason, in our opinion, to further investigate the properties of boson stars.  Boson stars with a repulsive self-interaction mediated by vector mesons within the effective theory framework were studied only recently \\cite{agni}. It is found that mass-radius curves satisfy the scaling properties for arbitrary boson masses and interaction strengths. There exist an another novel possibility of having a boson star with a composite scalar field. Such a composite scalar field can possibly arise when the superconducting colour quark matter undergo a transition from the BCS phase to Bose-Einstein condensation phase via a so called BCS-BEC transition \\cite{pir1, bec}. It should be noted that BCS regime of the colour superconducting phase is well studied using the techniques of perturbative QCD. But in BCS-BEC crossover can occur in low density and low temperature regime, where the coupling constant is large and the perturbative techniques are not available. One may use the methods of effective theory to describe the matter in the strongly coupled regime. In this picture, instead of a repulsive vector-meson exchange one may have repulsive self-interaction of the scalar field describing the (diquark) condensates \\cite{pir1}. A diquark-BEC phase may occur at low density regime somewhere between colour neutral nuclear matter and the free quark phases. This requires an additional constraint namely - the size of the condensates or the coherent length $\\xi$ should not be larger than the inter-particle spacing i.e. $n^{-1/3}$ , where $n$ is the number density.  In this work we analyze the equilibrium solutions of a boson star in a low temperature finite chemical potential regime. As far as we know boson stars  with a finite chemical potential are not properly considered in the literature. Majority of study do not focus on the statistical aspect of the self-gravitating boson gas but they rather appeal to find a specific equilibrium (quantum) solution of the Einstein-Klein-Gordon system. In the context of quark star we would like to mention that possibility of diquark star has already been considered in the literature \\cite{diquarkstar}. However the effect of Bose-Einstein condensation was not considered in this work. This effect can be significant in a low temperature finite chemical potential regime \\cite{pir1}. Moreover, in general relativity temperature and chemical potential are local functions of space and time and therefore they depend upon the metric. The metric dependence of these quantities can be specified by using the well-known Tolman conditions for a thermal equilibrium in the gravitational field \\cite{tol34, bilic99}. Again this aspect was not considered in the literature on boson or diquark stars. In this work we derive a generalized set of Tolman-Oppenheimer-Volkov equations by incorporating the effect of metric dependence of chemical potential. This has resulted in an additional differential equation describing the space-time evolution of the chemical potential coupled with the TOV \\cite{tov} equations.  In a very high density system such as a boson star, the introduction of finite chemical potential can rise a question. However in the standard model, for example, quark numbers are conserved and one can introduce chemical potential for them. In addition for a BCS-BEC crossover kind of transitions chemical potential play a very important role \\cite{pir1}. Chemical potentials are known to alter the global minimum of a scalar-field. In this work have demonstrated that when the scalar field mass $m$ is comparable to the chemical potential $\\mu$, the mass-radius relations and the maximum mass of the star are significantly changed.  In what follows we first derive the generalized set of TOV equations and study their numerical solutions by providing the equation of state as an input. The equation of state is derived in the mean-field approximation. We use for the parameter values of the diquark phase in a manner consistent with that given in Ref. \\cite{pir1}. ",
        "conclusions": "Before solving stellar structure equations numerically let us consider the following qualitative arguments. First consider the case when there is no self-interaction term and also no chemical-potential i.e. $g=0$ and $\\mu=0$. In this case the critical mass and radius of the star given $R\\sim \\frac{1}{m}$ and $M=M_{planck}^2/m$ \\cite{kaup}. Next, if the boson mass is around $100 GeV$, the corresponding critical mass of the star would be  $M=10^{17} M_{planck}\\sim 10^9 kg\\sim 10^{-21}M_N$ and its radius $R\\sim10^{-22} R_N$ where,$M_N$ and $R_N$ are mass and radius of a neutron star. Consequently, the  density of such a boson star would be $10^{45}$ times that of neutron star. Thus for a diquark-condensate with mass $(0.5-1)GeV$ \\cite{pir1}, the density of the corresponding star would be $10^{41}$ times of that of a neutron star. This is an extremely large value of the matter density for a diquark star to exist. However, things can change if the self-interaction term is nonzero. It is known that for the case $g\\neq 0$ (and $\\mu=0$), the scalar-field $\\mid\\Phi\\mid$ scales as $M_{Planck}/\\Lambda$ if $\\Lambda=\\frac{gM^2_{planck}}{2\\pi m^2} \\gg 1$\\cite{shapiro}. This condition is  to satisfy even for a very small values of $g$. Comparing the quadratic and quartic  terms in the potential $U$, we get $\\frac{g\\mid\\Phi\\mid^4}{m^2}\\sim \\Lambda\\frac{\\mid\\Phi\\mid^2}{M^2_{planck}}\\sim O(1) $. The energy density $\\rho$ will be  about $ m^2M^2_{planck}/\\Lambda$. Thus the energy density would corresponds to that of a non-interacting boson if mass is rescaled as $ m\\rightarrow m  \\sqrt{\\frac{1}{\\Lambda}}$. From this the maximum mass can be found \\cite{kaup, maxmass} to be \\begin{equation}\\label{mcrit} M_0\\approx \\frac{2}{\\pi}\\sqrt{{\\Lambda}} \\frac{M^2_{planck}}{m} \\end{equation} \\noindent This is also matches with the estimate critical mass obtained numerically in Ref. {\\cite{shapiro}. In our case, in the flat space, the classical potential has non-trivial minimum at $\\mid\\Phi\\mid^2=\\frac{\\mu^2-m^2}{2g}$. We have assumed that $m^2>0$ and for $\\mu^2>m^2$ we have condensation. In this case the mass scaling will be different and it is given by $m\\rightarrow m/\\sqrt{\\Lambda}\\left(\\frac{\\mu^2}{B}-1\\right)$ and the maximum mass may be given by $\\sqrt{\\Lambda}\\frac{M^2_{planck}}{m}\\left(\\frac{\\mu^2}{B}-1\\right)$. This suggest that the maximum mass should decrease if the chemical potential is decreasing. Since the chemical potential term here depend on space-time the above is only a rough estimate of the maximum mass. It also must be noted that the above arguments are not valid if $\\frac{(\\mu/m)^2}{B}<1$.  The validity condition $\\xi < n^{-1/3}$ for the condensate description. $\\xi(=\\frac{1}{\\Delta})$ can be found from the gap equation \\cite{raja} \\begin{equation} \\Delta = 2\\mu exp\\left( -\\frac{3\\pi^2}{\\sqrt{2}g}\\right), \\end{equation} \\noindent for $\\mu\\approx 400-500 MeV$ and the  corresponding running coupling constant $3.56$. However, at this juncture  the above equation should not be taken too seriously as it is based upon the arguments of the perturbation theory. The BCS-BEC transition is supposed to take place  in the high coupling regime. The quark-condensation condition in the flat-space is  ${(\\mu/m)^2}>1$ \\cite{pir1}. However, in the case of a boson star it needs to be locally satisfied.  Next,  we write the generalized TOV equations  in a dimensionless form. We have already introduced the dimensionless variables for the pressure $\\tilde{p}$, density $\\tilde{\\rho}$ and chemical potential $\\tilde{\\mu}$. We calculate the mass $M$ in the dimensionless unit $\\tilde{M} = M/M_o$ where $M_o$ is the solar mass. Radius of the star ($R$) is also computed in terms of $\\tilde{R} = R/R_o$, with $R_o = G M_o = 1.47 km$. Radius of the star can be defined as the distance from the center where $\\rho$ becomes zero. We solve equations (\\ref{tov2}, \\ref{tov1}, \\ref{tov3}) by specifying various values of the central pressure $p_c$ and the chemical potential $\\mu_c$, while $M$ at the center is considered to be zero.  From the form of equation (11) it is clear that, the energy density could be negative if the values $p_c$ and $\\mu_c$ are chosen arbitrarily. Therefore it is required to constrain their values by imposing the condition $\\tilde{\\rho} > 0$.\\\\   \\paragraph*{$\\tilde{\\rho} > 0$ condition and two branches of solutions}: \\begin{figure} \\includegraphics[width=7cm,height=7cm]{branches.eps} \\caption{The graph shows the solutions for initial conditions of pressure $\\tilde{p}_c ~\\text{and chemical potential} ~ \\tilde{\\mu}_c$ for $\\tilde{\\rho}_c\\geq 0$.}\\label{branches} \\end{figure}  The condition (on the initial conditions) implies that there are following two branches of the solutions:  \\begin{eqnarray} \\tilde{p}&\\geq& \\frac{16}{81} \\left(9-15 \\tilde{\\mu} ^2+8 \\tilde{\\mu} ^4\\right)\\\\ \\nonumber &-&\\frac{32}{81} \\sqrt{-27 \\tilde{\\mu} ^2+72 \\tilde{\\mu} ^4-60 \\tilde{\\mu} ^6+16 \\tilde{\\mu} ^8};~ (1<\\tilde{\\mu} \\leq \\sqrt{\\frac{3}{2}})\\\\ \\tilde{p}&\\geq& \\frac{16}{81} \\left(9-15 \\tilde{\\mu} ^2+8 \\tilde{\\mu} ^4\\right)\\\\ \\nonumber &+&\\frac{32}{81} \\sqrt{-27 \\tilde{\\mu} ^2+72 \\tilde{\\mu} ^4-60 \\tilde{\\mu} ^6+16 \\tilde{\\mu} ^8};~(\\tilde{\\mu} >\\sqrt{\\frac{3}{2}}) \\end{eqnarray}   The solid curve in figure (1) represents the first branch given by equation (25), which is valid the values of chemical  potential defined in the range $1 <\\tilde{\\mu} \\leq \\sqrt{\\frac{3}{2}}$. The curve with the dotted line represents the branch given by equation (26). From figure (1) it is clear that solutions of the equation (25) is valid even in the region $1 <\\tilde{\\mu} \\leq \\sqrt{\\frac{3}{2}}$.   \\begin{figure} \\includegraphics[width=7cm,height=7cm]{Nodes.eps} \\caption{For a given chemical potential the density profile can have nodes. The upper most curve is obtained by imposing the no-node condition.} \\end{figure}  Figure (2) shows the density profile as function of $\\tilde{r}$ for various values of the central pressure and the chemical potential $\\tilde{\\mu}=1.5$. The solution indicates the nodes in the profiles. However, this is a unphysical behavior and it  can be avoided by imposing the no-node condition on $\\rho$.\\\\   \\paragraph*{No-Node condition}: We need to make sure that $\\tilde{\\rho}'(r=0) < 0$ inside the star. Analytically this condition can be represented as, \\begin{eqnarray} \\left(2+3 \\sqrt{\\tilde{p}(0)}\\right) \\left(\\sqrt{1+\\sqrt{\\tilde{p}(0)}}-\\tilde{\\mu }(0)\\right) \\tilde{p}'(0)\\\\ \\nonumber -4 \\left(1+\\sqrt{\\tilde{p}(0)}\\right) \\tilde{p}(0) \\tilde{\\mu }'(0)<0 \\end{eqnarray}  The upper most curve in figure (2) is plotted by incorporating the no-node condition on the initial values of pressure and chemical potential. Thus the implementation of the no-node condition shows that for a given chemical potential there exist a lower bound on the central pressure. If one specify the central pressure below this bound then the nodes develops in the density profiles. Our numerical study indicates that if the initial values of the pressure and chemical potential are chosen just above the two branches of '$\\tilde{\\rho}> 0$ condition' as shown in figure (1), the density profile develops the nodes. One can avoid the nodes for higher values of the central pressure $p_c$, but this will also increase the central density as implied by equation (\\ref{eos}). It should be noted that the nodes in our analysis arise solely due to the introduction of the chemical potential and therefore their origin is different than the nodes in the wave functions reported in earlier literature of boson-star (for example \\cite{shapiro}).   \\begin{figure} \\includegraphics[width=7cm,height=7cm]{mass-p-mu.eps} \\caption{Maximum Mass of the star versus initial pressure for different values of $\\tilde{\\mu}_c$ is plotted. The peak of the graph in each case denotes the highest maximum mass configuration.} \\end{figure}  Figure (3) shows graph of  the star-mass as a function  of the central pressure for  different values of the chemical potential. The lower most curve represents the case of zero chemical-potential. The mass is increasing with increasing the central pressure (density) and it reaches a plateau. This behavior is in agreement with the earlier work. Also the value of maximum-mass is in agreement with Ref. \\cite{shapiro}. The introduction of a finite chemical-potential increases the overall value of the maximum-mass and the plateau. This figure shows that when the values of  $\\mu_c$, increases, the maximum-mass values occurs at higher values of $p_c$. The lower values of $p_c$ in this case ($\\mu_c>1$) are ruled out due to the no-node condition. If this condition is violated then the maximum-mass configurations would occur at much lower $p_c$. And in this case the  maximum-mass would be much higher than the $\\mu_c=1$ case. Thus the consistent solutions with finite  chemical-potential requires to have higher values of the central-pressure in comparison with the $\\mu_c=0$ case.   \\begin{figure} \\includegraphics[width=7cm,height=7cm]{M-R-curves.eps} \\caption{Maximum Mass of the star versus Radius for different values of $\\tilde{\\mu}_c$ is plotted. The peak of the graph in each case denotes the highest maximum mass configuration.}\\label{MR} \\end{figure}  In figure (4) we have plotted mass of the star versus radius for different values of the chemical-potential. The radius is determined by the value of radial coordinate $r~(=R)$, where the density vanishes. The metric-dependence of the chemical-potential gives a new class of solutions. Figure (4a) shows the case of zero central chemical-potential, in this case there is no variation in the chemical-potential as implied by equation (\\ref{tov3}). In this case, the maximum mass value occurs at the peak of the curve and its value matches well with the known result \\cite{shapiro}.  All the points on the left-hand side of the peak are unstable, as they have higher gravitational potential energy than the points on the right-hand side. Figure (4b) represents  $\\mu_c=1$ case, one can see that the value of maximum-mass has increased as compared to $\\mu_c=0$ case and more strikingly the value of the maximum-mass does not occur on the peak of the curve.  The spiral like structure in the lower part of the curve represents two branches of then solution. In this region for any given $R$ there exist two points with the same mass. The points on the lower-arm of the spiral have minimum energy, while points on the upper-arm may be unstable. Figure (4c) shows the case with $\\mu_c=1.2$. Again the maximum-mass does not occur on the peak of the curve and its value is smaller than $\\mu_c=1$ case. Radius of the star at the maximum-mass point is much smaller than the case shown in figure (4b) and thus the density of the star is higher than $\\mu_c=1$ case. Figure (4d) shows that for $\\mu_c=1.5$, the maximum-mass is smaller than $\\mu_c=1.2$ case and it occurs at a smaller radius. In this case the maximum-mass point lie in the unstable region. The lower-arm of the spiral represents the stable branch of the solution. The maximum-mass point in the stable branch  is comparable to the case shown in figure (4a). This points occurs when $(\\frac{d\\tilde{M}}{d\\tilde{R}})^{-1}=0 $ is satisfied. It should be noted that the qualitative arguments for maximum-mass of a star, given below equation (22),  may not be applicable for the cases shown in figure (4b-4d). As these arguments do not account for the variation in the chemical-potential and no-node condition imposed on the $\\rho$.    Figures (5-6) show  pressure and density as functions of $\\tilde{r}$ for the case when the central-pressure is kept fixed at $\\tilde{p}_c=35$ and vary $\\tilde{\\mu}_c$. Figure (5) indicates that the pressure decreases less slowly with $r$ by increasing the value of chemical-potential. Thus the star radius may increase if the chemical-potential is increasing. The density vs. distance plot also shows the similar behavior. In this case the value of central-density $\\tilde{\\rho}_c$ is not fixed, but it decreases as the chemical-potential at center increases. This behavior is expected from the form of the EOS (\\ref{eos}). \\begin{figure} \\includegraphics[width=7cm,height=7cm]{p-r-curves.eps} \\caption{ The variation of the normalized pressure with the normalized radius for the different values of the central chemical-potential is shown here. The central pressure in all the cases is kept fixed.} \\end{figure} \\begin{figure} \\includegraphics[width=7cm,height=7cm]{rho-r-curves.eps} \\caption{ The variation of the normalized density with normalized radius for the different values of the central-chemical potentials is shown here. The central pressure in all the cases is kept fixed. } \\end{figure}  Figure (7) shows the typical variation in  chemical-potential and the metric-function $B$ with the radius. The metric-function approaches the unity at the boundary, while the value of the chemical-potential decreases from its central value with the radius. However, the value of $\\tilde{\\mu}$ remains non-zero at the boundary. The dynamics of the chemical-potential and $B$ are complementary as can be anticipated from definition of $\\tilde{\\mu}$. \\begin{figure} \\includegraphics[width=7cm,height=7cm]{muB-r.eps} \\caption{The typical variation of the normalized chemical-potential and the metric-function $B$ with the radius is shown for the given values of central pressure and chemical-potential.} \\end{figure}   Let us consider the case when the scalar-field represents the diquark-condensates which may arise in a low-density and strong-coupling regime. The condition for the condensate formation is $(\\mu/m)^2/B>1$ or $\\tilde{\\mu}^2>1$. One can notice from figure (7) that this condition is violated somewhere inside the star. For the condition to be valid inside the entire region of the star one has to increase the value of $\\tilde{\\mu}_c$. From the discussion above, increasing $\\tilde{\\mu}_c$ would also requires to increase the central-pressure $\\tilde{p}_c$. This may in turn increases the density as shown in figure (2). Clearly the energy density can not increase to a very high values otherwise the condition of BEC formation may not be fulfilled. For example one may take the value of the chemical potential $\\tilde{\\mu}_c=1.5$ which is consistent with \\cite{pir1}. According to the figure (2), the physical density profile should be around 30 times larger then the nuclear density which is not consistent with the BCS-BEC crossover assumption. Moreover, the mass-radius diagrams shown in figure (4) suggest that as $\\tilde{\\mu}$ increases the available parameter-space for the stable configuration reduces significantly. We would like to add here that when one considers the case of fundamental scalar-field, the constraint of the low energy-density is not required. In this situation the boson-star equilibria are defined over a large range of the parameters."
    },
    "0910/0910.1834_arXiv.txt": {
        "abstract": "We introduce a data reduction package written in Interactive Data Language (IDL) for the Magellan Echellete Spectrograph (MAGE).  MAGE is a medium--resolution ($R \\sim$ 4100), cross--dispersed, optical spectrograph, with coverage from $\\sim$ 3000 -- 10000 \\AA.   The MAGE Spectral Extractor (MASE) incorporates the entire image reduction and calibration process, including bias subtraction, flat fielding, wavelength calibration, sky subtraction, object extraction and flux calibration of point sources.   We include examples of the user interface and reduced spectra.  We show that the wavelength calibration is sufficient to achieve $\\sim 5$ km s$^{-1}$ RMS accuracy and relative flux calibrations better than $10\\%$.  A light-weight version of the full reduction pipeline has been included for real--time source extraction and signal--to--noise estimation at the telescope. ",
        "introduction": "Cross--dispersed spectrographs are attractive instruments for a variety of science applications, primarily for their broad wavelength coverage and high spectral resolution.   While these instruments maximize wavelength coverage and spectral resolution, the reduction and calibration of their data is complex.  The spatial and dispersion axes are typically not aligned with the CCD columns and rows; a given wavelength can span several columns on the chip.  Furthermore, the curvature and tilting of the orders complicates the tracing of order edges and the extraction of spectra.  Fortunately, many of these issues have been addressed by modern reduction techniques \\citep[e.g.,][]{1982A&A...111..260M,1985A&A...143...13R,1986SPIE..627..707P,1988PASP..100.1572G, 1989PASP..101.1032M, 1990PASP..102..183M, 1994PASP..106..315H, 2002A&A...385.1095P, 2003PASP..115..688K}.  Pipelines exist and are regularly employed for extracting cross--dispersed optical \\citep{2002A&A...385.1095P} and infrared \\citep{2004PASP..116..362C} data.  Large, multi--fiber spectroscopic surveys, such as the Sloan Digital Sky Survey \\citep[SDSS;][]{2000AJ....120.1579Y, 2002AJ....123..485S, 2002SPIE.4847..452S}, the Large Sky Area Multi--Object Fiber Spectroscopic Telescope \\citep[LAMOST;][]{2008AcASn..49..327C} and the Radial Velocity Experiment \\citep[RAVE;][]{2006AJ....132.1645S} also use pipeline reduction techniques that are directly applicable to the analysis of cross--dispersed data.  We have developed a new reduction package for the Magellan Echellete spectrograph \\citep[MAGE;][]{2008SPIE.7014E.169M}.  The MAGE Spectral Extractor (MASE) is written in the Interactive Data Language (IDL) and is based on the SDSS spectro2d pipeline\\footnote{Available at \\url{http://spectro.princeton.edu/}}, the Magellan Inamori Kyocera Echelle \\citep[MIKE;][]{2003SPIE.4841.1694B} IDL pipeline\\footnote{Available at http://web.mit.edu/\\~{}burles/www/MIKE/} (Bernstein, Burles \\& Prochaska, in prep) and the XIDL astronomy package\\footnote{Available at http://www.ucolick.org/\\~{}xavier/IDL/index.html}.  The software employs a graphical user interface (GUI) built with IDL widgets to control the reduction process.  A faster, ``quicklook\" version of the pipeline is packaged with MASE, for use at the telescope.   This version can be used for on--the--fly reductions and provides estimates of signal--to--noise ratios and spectral morphology.  In this paper, we describe the details of the MASE reduction pipeline and GUI.  In \\S \\ref{sec:mage}, we review the properties of the MAGE spectrograph.  The reduction pipeline is detailed in \\S \\ref{sec:mase}, with examples of the MASE GUI and reduced spectra.  Our conclusions are reported in \\S \\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary} We have developed a new data reduction package for the MAGE spectrograph.  The MASE GUI and pipeline is written in IDL and incorporates modern sky--subtraction and optimal extraction techniques.   The program is available for download online.  The pipeline is designed for reductions of unresolved point sources.  However, for spectra of extended objects, MASE can still be used to bias-subtract, flat field and wavelength calibrate the observations.  Future versions of MASE will include a telluric correction routine, which will require additional calibration observations, and scripting capabilities for reducing multiple nights of data.  Support for the design and construction of the Magellan Echellette Specrograph was received from the Observatories of the Carnegie Institution of Washington, the School of Science of the Massachusetts Institute of Technology, and the National Science Foundation in the form of a collaborative Major Research Instrument grant to Carnegie and MIT (AST--0215989).   J.X.P. is partially supported by an NSF CAREER grant (AST--0548180).   We thank the referee for their constructive comments that improved this manuscript.  J.J.B. acknowledges Jackie Faherty and Dagny Looper for their illuminating conversations and extensive vetting of MAGE data. We thank Ricardo Covvarubias for early testing of the MAGE quicklook tool.  Finally, thanks to the entire Magellan staff, in particular Mauricio Martinez and Hern\\'{a}n Nu\\~{n}ez."
    },
    "0910/0910.4730_arXiv.txt": {
        "abstract": "{ We make use of three dimensional hydrodynamical simulations to investigate the effects of granulation on the \\ion{Cu}{i}  lines of Mult.\\,1 \\relax  in the near UV, at 324.7\\,nm and 327.3\\,nm. These lines remain strong even at very low metallicity and provide the opportunity to study the chemical evolution of Cu in the metal-poor populations. We find very strong granulation effects on these lines. In terms of abundances the neglect of such effects can lead to an overestimate of the A(Cu) by as much as 0.8 dex in dwarf stars. Comparison of our computations with stars in the metal-poor Globular Clusters NGC 6752 and NGC 6397, show that there is a systematic discrepancy between the copper abundances derived from Mult.\\,  \\relax 2 in TO stars and those derived in giant stars of the same cluster from the lines of Mult.\\,2 at at 510.5\\,nm and 587.2\\,nm. We conclude that the \\ion{Cu}{i} resonance lines are not reliable indicators of Cu abundance and we believe that an investigations of departures from LTE is mandatory to make use of these lines. ",
        "introduction": "In our quest to understand how the Universe evolved from the primordial chemical composition, consisting of H, He, and traces of Li, to the present day complexity, we strive to uncover all the nucleosynthetic channels. This requires to measure the evolution of as many chemical species as possible. For some species this becomes difficult at low metallicities, when all the observable lines become very weak. This is the case for Cu, the measurements of Cu abundances are mainly based on the lines of Mult. 2 at 510.5\\,nm and 578.2\\,nm \\citep{Sneden91,Mishenina}. However, such lines become very weak at low metallicities, even in cool giants. The strongest line, at  510.5\\,nm, has an equivalent with of a few tenths of picometer for a K giant of metallicity --2.5 and becomes very difficult to measure. This induced several groups \\citep{Bihain,Cohen08} to push the observations in the near UV, where the \\ion{Cu}{i} resonance lines at 324.7\\,nm and 327.3\\,nm are stronger by  a factor of ten and can be measured at the lowest metallicities. These lines are formed in the cool outer layers of the stellar atmosphere. Such layers are formally stable against convection and 1D model atmospheres cannot account for the effects of convective motions (``granulation'') here. The use of 3D hydrodynamical simulations has shown that one effect of this is a steeper temperature gradient in the outer layers than what predicted by 1D models \\citep{A99,A05}. This effect is often referred to as ``overcooling'' and is more pronounced for metal-poor stars. In this contribution we wish to investigate the effects of granulation on the formation of the \\ion{Cu}{i} resonance lines in metal-poor stars. To do so we analyse the lines in the turn-off stars of two metal-poor Globular Clusters (NGC 6752 and NGC 6397) and compare the derived abundances, both in 1D and 3D, with those derived for giants, making use of the lines of Mult.2 ",
        "conclusions": "From Table \\ref{abun} it is immediately clear that the 3D corrections are large. The reasons for this can be understood by looking at Fig.\\,1 where the temperature distribution of one of our 3D models is depicted, together with the mean temperature distribution and that of a corresponding LHD model. The overcooling is obvious from the top panel and the contribution functions in the two bottom panels reflect  the fact that in these cool outer layers the Cu atoms populate mainly the ground layer, thus contributing for the bulk of the absorption of the line, at variance with what happens in the 1D model (dashed line). The two bottom panels correspond to two different Cu abundances, illustrating how the tendency to prefer the outer layers increases with the increasing number of absorbers. As expected this results in larger 3D corrections for more metal-rich stars. However, one should be always cautious when facing contribution functions like the ones shown in Fig.\\,1. In fact all the computations have been performed in LTE, it is likely that photons coming from the warm streams may produce overionisation in these low density outer layers. If NLTE effects are important a considerable resizing of the outer peak of the contribution function can be expected. A way to check indications of possible NLTE effects is to compare the abundances in the cluster, derived from the resonance lines in TO, with those of giants, derived from the lines of Mult.\\,2. We derived  abundance in NGC\\,6397 using the EWs for two giants measured by \\citet{gratton82}, for NGC\\,6752 we used a UVES spectrum of a giant star (star  Cl* NGC 6752 YGN 30), already studied by \\citet{Yong}. Neither in 1D nor in 3D giants and dwarfs provide the same abundance. In NGC 6752 the dwarf stars imply a higher abundance, both in 1D and 3D, the reverse is true in NGC 6397. That the problem is with the modelling of the Cu lines is confirmed by the analysis of the giant in NGC\\,6752, for which we were able to measure both the \\ion{Cu}{i} resonance line at 327.3\\, nm and the 510.5\\,nm line. Both in 1D and 3D the two lines provide abundances which differ by about 0.5\\,dex. We conclude that the  \\ion{Cu}{i} resonance lines are not good abundance indicators, a full 3D-NLTE study should be undertaken for these lines. At the same time one may  suspect that also the lines of Mult.\\,2 may be affected by deviations from LTE. The Galactic evolution of copper must be placed on solid grounds with a better modelling of the line formation.     \\balance"
    },
    "0910/0910.4395_arXiv.txt": {
        "abstract": "Eta Carinae shows broad peaks in near-infrared (IR) $JHKL$ photometry, roughly correlated with times of periastron passage in the eccentric binary system.  After correcting for secular changes attributed to reduced extinction from the thinning Homunculus Nebula, these peaks have IR spectral energy distributions (SEDs) consistent with emission from hot dust at 1400--1700~K.  The excess SEDs are clearly inconsistent, however, with the excess being entirely due to free-free wind or photospheric emission.  One must conclude, therefore, that the broad near-IR peaks associated with Eta Carinae's 5.5 yr variability are due to thermal emission from hot dust.  I propose that this transient hot dust results from episodic formation of grains within compressed post-shock zones of the colliding winds, analogous to the episodic dust formation in Wolf-Rayet binary systems like WR~140 or the post-shock dust formation seen in some supernovae like SN~2006jc.  This dust formation in Eta Carinae seems to occur preferentially near and after periastron passage; near-IR excess emission then fades as the new dust disperses and cools.  With the high grain temperatures and Eta Car's C-poor abundances, the grains are probably composed of corundum or similar species that condense at high temperatures, rather than silicates or graphite.  Episodic dust formation in Eta Car's colliding winds significantly impacts our understanding of the system, and several observable consequences are discussed. ",
        "introduction": "As one of the brightest mid-infrared (IR) sources in the sky, Eta Carinae has been a key reservoir of information about dust around massive stars (Chesneau et al.\\ 2005; Smith et al.\\ 1998, 2003b; Aitkin et al.\\ 1995; Robinson et al.\\ 1987; Hackwell et al.\\ 1986). Dust formation by massive stars is of considerable interest regarding the presence of dust in star-forming galaxies at high redshift (e.g., Bertoldi et al.\\ 2003), when the Universe was too young for asymptotic giant branch stars to have contributed significantly (Tielens 1998). The neutral and dusty bipolar Homunculus Nebula, ejected in the 19th century Great Eruption, absorbs most of the ultraviolet and visual-wavelength radiation from the star and re-radiates that luminosity in the mid-IR, providing a handy calorimeter of the central star (Smith et al.\\ 2003b).  This paper, however, focusses attention not on the dusty Homunculus, but on the near-IR emission from the very core of the nebula surrounding the central star (Rigaut \\& Gehring 1995; Smith \\& Gehrz 2000; Chesneau et al.\\ 2005), which exhibits distinct and peculiar variability in the near-IR.  Whitelock et al. (1983, 1994, 2004) have followed the photometric near-IR variability of Eta Carinae, and noted quasi-periodic peaks every $\\sim$5 yr.  Damineli (1996) proposed that this IR variability and Eta Car's optical spectroscopic changes were due to an eccentric 5.5 yr binary system, and successfully predicted the time of the next periastron passage in 1998.0.  Since then, an intensive multiwavelength campaign has documented Eta Car's cyclical variability in the X-rays (Corcoran 2005; Corcoran et al.\\ 2001; Ishibashi et al.\\ 1999), optical (Damineli et al.\\ 1998; Davidson et al.\\ 2000), and radio (Duncan \\& White 2003).  Here we return to the near-IR emission peaks that first heralded the $\\sim$5 yr variability, which have so far eluded a satisfying explanation.  We are interested in the broad $\\Delta t \\simeq$1--2 yr peaks in the $JHKL$ light curves, and not the narrow eclipse-like events in the month or so surrounding periastron (Whitelock et al.\\ 2004; Feast et al.\\ 2001), which probably have a different origin and will be the topic of a future paper.  In particular, $JHKL$ photometry obtained from 1972 to 2004 by Whitelock et al.\\ (1983, 1994, 2004) is used to constrain the spectral energy distribution (SED) of the {\\it excess} emission that rises above the more steady, secular increase in apparent brightness over recent decades noted by Whitelock et al.\\ (1994).  We show that the broad quasi-periodic peaks in the $JHKL$ light curves are caused by emission from hot dust grains, and we propose that these grains condense within the dense post-shock gas in the colliding-wind binary system.  SEDs of this excess emission are not consistent with photospheric or free-free emission, which was previously assumed to be the origin.   Section 2 reviews the observational data, introducing a correction for the slow increase in apparent brightness caused by the expanding and thinning Homunculus Nebula.  Section 3 presents the excess-emission SEDs for several of the past periastron passage events, and discusses possible emission mechanisms and the properties of the dust responsible for the excess near-IR emission.  Finally, Section 4 proposes a scenario wherein new dust grains form in post-shock gas around times of periastron in Eta Carinae as a result of compression by the colliding winds.    \\begin{figure}\\begin{center} \\includegraphics[width=3.1in]{fig1.eps} \\end{center} \\caption{Near-IR $JHKL$ photometry of Eta Car (solid points) converted to flux, from Whitelock et al.\\ (2004).  The solid curve shows the adopted secular trend of the increasing continuum level, probably due to the expansion and thinning of the Homunculus Nebula and the consequent decrease in extinction to the central star.  The unfilled points show the photometry after dividing by this secular trend.}\\label{fig:plotJHKL} \\end{figure} \\begin{figure}\\begin{center} \\includegraphics[width=3.0in]{fig2.eps} \\end{center} \\caption{Near-IR extinction law for the Homunculus of $\\eta$ Car (solid points) normalized to the $L$-band extinction $A_L$, determined from the solid curves in Figure~\\ref{fig:plotJHKL} that approximate the secular brightening trend in $JHKL$ photometry.  The dashed line shows the average ISM extinction law in the Milky Way from Indebetouw et al.\\ (2005) for comparison.}\\label{fig:extinction} \\end{figure} \\begin{figure}\\begin{center} \\includegraphics[width=3.1in]{fig3.eps} \\end{center} \\caption{Similar to Figure~\\ref{fig:plotJHKL}, but the solid points show the near-IR flux with a correction for the slowly-changing extinction that is presumed to cause the secular brightening over this time period.  Total extinction at the latest epochs is assumed to be negligible.  The unfilled points show the $JHKL$ excess flux, with the constant baseline flux (due to stellar photospheric and wind emission) subtracted, leaving only the excess near-IR flux associated with the 5.5 yr variability. The short horizontal dashes show the estimated flux for each dust formation event.}\\label{fig:plotJHKLsub} \\end{figure} ",
        "conclusions": "\\subsection{Observed Precedents for Post-shock Dust Formation}  Although one may expect hot, X-ray emitting post-shock gas to be hostile to the formation of grains that require temperatures below 2000~K, two different classes of object have revealed strong evidence for the condensation of new dust grains in fast shocks.  Both classes, discussed below, share interesting similarities with Eta Carinae.  First, some Wolf-Rayet (WR) binaries of WC type show clear evidence for dust formation in their winds (Gehrz \\& Hackwell 1974).  Perhaps the most relevant example is the episodic dust formation in the eccentric WC binary WR~140 (Monnier et al.\\ 2002; Williams et al.\\ 1990; Hackwell et al.\\ 1979), whose colliding winds give rise to an X-ray light curve that is similar in several respects to that of Eta Car (Pittard \\& Dougherty 2006).  WR~137 and WR~19 are similar examples (Williams et al.\\ 2001, 2009).  There are also several cases of WC binaries with more circular orbits that produce the so-called ``pinwheel'' nebulae.  High resolution imaging of these systems (Tuthill et al.\\ 1999; Monnier et al.\\ 1999) provides strong evidence that their IR emission comes from dust forming in post-shock gas as a consequence of their colliding winds.  Dust may form in the cone of post-shock gas as it flows away from the central stars and cools below condensation temperatures.  The C-rich composition of the winds aids the efficiency of amorphous carbon dust formation (Gehrz \\& Hackwell 1974).  Second, recent studies of some supernovae (SNe) that interact with dense circumstellar matter also provide clear evidence for the condensation of new dust grains in the post-shock gas.  The best case is SN~2006jc, which at only $\\sim$50 days after explosion, simultaneously showed a fading of its visual light due to extinction, an increase in its IR flux caused by dust at $\\sim$1600~K, and a systematic blueshift of its emission lines as the far side of the SN was blocked by the new dust.  Smith et al.\\ (2008a) first demonstrated these changes and showed that the newly formed dust must be in the post-shock gas.  SN~2006jc is particularly relevant to Eta Car, because at the same time that SN~2006jc formed dust and exhibited an IR excess, it also showed a boost in its X-ray luminosity (Immler et al.\\ 2008) and enhanced He~{\\sc ii} $\\lambda$4686 emission (Smith et al.\\ 2008a).  He~{\\sc ii} $\\lambda$4686 in Eta Car is also associated with the periastron passages when dust forms (Steiner \\& Damineli 2004; Martin et al.\\ 2006), as is enhanced X-ray emission.  As noted by Smith et al.\\ (2008a), these similarities implicate common physical conditions in the post-shock cooling zones of Eta Car and SN~2006jc.  While both WC binaries and SN~2006jc have C-rich material that may aid the formation of graphite grains, this is not necessarily a prerequisit for dust formation in post-shock gas.  There is now a growing number of additional cases for post-shock dust formation in Type IIn supernovae that are {\\it not} expected to have C-rich ejecta (Smith et al.\\ 2009, 2008b; Fox et al.\\ 2009; Pozzo et al.\\ 2004). Details of how dust forms in the hot post-shock gas is not yet fully understood, but these observed examples prove that Nature can overcome the difficulties involved.  Strong radiative cooling in dense post-shock gas is probably instrumental (Usov 1991), as are density inhomogeneities.    \\subsection{Post-Shock Grain Condensation in the Colliding Winds of Eta Car}  Given the episodic presence of the hot dust in Eta Car, the most promising place for dust to form is in the compressed post-shock primary wind, akin to the episodic dust formation in WR~140 and other colliding-wind binary systems discussed above.  Recent hydrodynamic simulations of the colliding winds in Eta Car (Parkin et al.\\ 2009; Okazaki et al.\\ 2008) show dense gas clumps arising from instabilities at the colliding-wind shock, with $n_H > 10^{12}$ cm$^{-3}$ at radii of 2--3 AU around times of periastron.  These clumps will cool and become somewhat less dense as they flow down the shock cone, eventually reaching a radius where the radiative equilibrium temperature drops low enough for grains to condense.  For sublimation temperatures around $T\\simeq$1700~K, this occurs at $R \\ \\simeq \\ 60$~AU in Eta Car, which is only 2--3 times the orbital separation at apastron, and for which the flow timescale is only a few months at $\\sim$500 km s$^{-1}$.  Densities of clumps at these radii may drop to $\\sim$10$^9$ cm$^{-3}$, comparable to minimum densities for grain nucleation (e.g., Smith et al.\\ 2008a and references therein).  Grains may form at even smaller radii if high optical depths shield clumps from the primary star's radiation.  Once the dust grains have formed, they will cool quickly as they expand down the shock cone.  Flowing outward at 500 km s$^{-1}$ from an initial formation radius at 60 AU, for example, the dust equilibrium temperature will fall to $\\sim$1400 K after 100 days and will drop to $\\sim$1000 K after 1 yr.  This roughly matches the dust cooling indicated by the sequence of SEDs observed in the year following the 1987 periaston event (Fig.~\\ref{fig:sed1}).  The observed dust temperature does not necessarily need to drop however, because new dust may continue to form at the same temperature in the shock cone, replacing the dust that has cooled.  Some other events have SEDs showing relatively constant dust temperatures.  Long after the periaston event when the new dust has cooled, it should reside primarily in the complex expanding shells and loop structures created during the 5.5 yr orbit, like those seen in simulations of Eta Car's colliding winds (Parkin et al.\\ 2009; Okazaki et al.\\ 2008).  It would be quite interesting to make a detailed comparison between the geometry of these simulations and dust structures observed in the core of the Homunculus.  High-resolution data show an intriguing asymmetric distribution of material around the star (e.g., Smith et al.\\ 2004; Chesneau et al.\\ 2005; Gull et al.\\ 2009), in which high grain temperatures of 400--500~K are observed (Smith et al.\\ 1998, 2003b; Hackwell et al.\\ 1986).  Some evidence suggests that Eta Car's wind mass-loss rate changes from one cycle to the next (Damineli et al.\\ 1998, 2008; Davidson et al.\\ 2005), so the relevant post-shock density may be different in each cycle.  Depending on how efficiently the gas cools, when dust starts to form, and how long the dust formation episode lasts, significantly different IR excess emission may be observed in each successive periastron event, consistent with the somewhat irregular character of the $JHKL$ light curves.  With peak $\\lambda F_{\\lambda}$ values of (1--3)$\\times$10$^{-7}$ erg s$^{-1}$ cm$^{-2}$, the luminosity of the hottest dust that dominates the excess $JHK$ fluxes in Figures~\\ref{fig:sed1} and \\ref{fig:sed2} is roughly (4--12)$\\times$10$^4$ $L_{\\odot}$ (i.e. a small fraction of the star's radiative energy budget, but far exceeding the X-ray luminosity).  At temperatures of roughly 1400--1700~K, the dust mass needed to account for the IR luminosity is at least (1.4--2.8)$\\times$10$^{-8}$ $M_{\\odot}$ (adopting the same simplifying assumptions and methods as Smith \\& Gehrz 2005).  This is a lower bound to the mass formed in a periastron event because cooler dust is also present, as is evident from the SED shapes --- the hottest dust cools quickly as it flows downstream in the shock cone and may be replaced by additional hot dust (see Smith et al.\\ 2008a).  This is less than 0.1\\% of the total gas mass entering the shock interaction cone during the year around periastron.  Both the dust mass and the fraction of wind mass converted to dust in one of Eta Car's periastron events are quite similar to values estimated for WR~140, which are 2.8$\\times$10$^{-8}$ $M_{\\odot}$ and 0.2\\%, respectively (Williams et al.\\ 1990).  Broad peaks in the multi-filter IR light curves of WR~140 (Williams et al.\\ 1990) are also similar to those of Eta Car.  In dusty WC binaries and SN~2006jc, gas entering the shock is expected to be C enriched, favoring the formation of amorphous carbon grains. In Figures~\\ref{fig:sed1} and \\ref{fig:sed2}, we have compared the observed SEDs to graphite and silicates, as well as a Planck function with a more generic $\\lambda^{-1}$ emissivity that may be representative of some other grain species.  There are problems with attributing the new dust to either graphite or silicate grains. First, the high temperatures implied by the SED shapes are inconsistent with silicates and Fe-dominated grains, which seem to condense at lower temperatures of 1000--1200~K in novae (e.g., Gehrz 1988).  While some silicates may form and may be partly responsible for the cooler dust excess seen in the $L$-band, silicates cannot dominate the hottest $\\sim$1700~K dust seen at $JHK$.  The high temperatures are in principle consistent with graphite, which can condense at temperatures around 1600~K, but that would be surprising given the N-rich but severely O/C-poor abundances observed in Eta Car (Davidson et al.\\ 1986).  The dust mass formed in a periastron event is small compared to the available gas mass, so while C should be underabundant, the formation of some small amount of amorphous carbon is not necessarily precluded.  A more straightforward solution, however, may be to attribute most of the newly formed hot dust to another species, such as corundum (Al$_2$O$_3$), which condenses at high temperatures around 1700~K. Indeed, Chesneau et al.\\ (2005) argued that some contribution from corundum was needed to fit the very wide mid-IR ``silicate'' feature in Eta Car, and they found some spatial variation in the relative corundum abundance in the core of the Homunculus.  In particular, Chesneau et al.\\ (2005) found that a clump of dust located 0$\\farcs$4 southeast of the central star had a higher relative abundance of corundum than dust features in other directions, which they attributed to dust formation in a polar wind.  Smith (2006) showed that this SE clump was redshifted and therefore located on the far side of the star in the equatorial plane.  The fact that this equatorial clump is enhanced in corundum is quite interesting in the current context, because the orbit orientation favored by Okazaki et al.\\ (2008) places this clump in the direction of the secondary star at periastron passages.  In other words, if corundum dust forms preferentially around times of periaston in connection with the IR peaks in the photometry, and it forms in the colliding wind shock cone, then {\\it this is the preferential direction in which we would expect the new dust to flow}.  In WR~140, for example, puffs of episodic dust formed in the shock cone at periaston expand away from the star in that direction in high spatial resolution images (Monnier et al.\\ 2002).  This hints that further study of the spatial structure and composition of the dust in the core of Eta Car may prove to be extremely interesting.    \\subsection{Summary and Further Implications}  Following the precendent set by dusty WC binaries that form dust in their colliding-wind shocks and by some SNe that form dust in post-shock gas, this paper proposes that the broad peaks in the $JHKL$ light curve of Eta Car are due to the post-shock condensation of new dust in the colliding winds.  The dust forms around the time of periastron passage in the eccentric 5.5 yr binary orbit, probably within dense clumps that have expanded far enough from the star that the radiative equlibrium temperature drops below $\\sim$1700~K, and the dust then cools as it flows down the shock cone.  The clearest episodes were associated with with the 1981.4 and 1986.9 periastron events, although later events also show evidence for dust.  Episodic formation of dust in the colliding winds of the Eta Car binary system has several additional implications for this complex system:  1) Dust formed recently in this way is near the central star (within 60--500 AU), compared to the dust in the much larger Homunculus Nebula.  While it is a relatively small mass of dust, its proximity may cause substantial extinction.  This dust may preferentially block our view of the star without obscuring features in the surrounding nebulosity and extended wind, perhaps explaining Eta Car's mysterious ``coronograph'' described by Hillier et al.\\ (2006).  2) Absorption and scattering by this localized dust may be highly time variable as the dust expands and thins in the time interval between periastron passages.  This dust may block starlight, or may scatter starlight toward us, with different contributions at different times. It may therefore contribute to irregular variability in optical photometry of Eta Carinae.  Sterken et al.\\ (1998) noted a drastic reddening of visual colors after several events, which they attributed to (usually slower) S~Doradus variations, but perhaps the sudden reddening was from new dust close to the star instead.  3) Combined effects of local scattering and extinction may also influence spectra, even on very small angular scales of 0$\\farcs$1--0$\\farcs$2.  The asymmetric expanding dust formed in the colliding-wind shock cone could serve as a moving mirror, reflecting some specific lines of sight more than others.  This may be a factor in perplexing subtleties of spatially-dependent emission line profiles along our line of sight (e.g., Gull et al.\\ 2009; Davidson et al.\\ 2005), and direction-dependent profiles reflected by the Homunculus (e.g., Smith et al.\\ 2003a).  Even line profiles in our direct view of the central star may be affected by time-dependent dust scattering and extinction.  4) Episodic dust formation at small radii may regulate the escape of UV photons that excite emission lines in nearby gas condensations, contributing to the observed nebular spectral variability during the 5.5 yr cycle, and mimicing a shell ejection as described by Zanella et al.\\ (1984).  5) Dust production in successive events every 5.5 yr is not perfectly repeatable, signaling differences in the post-shock densities that may result from longer-term changes in the mass-loss rate of the primary star.  This, in turn, may alter the effects discussed above from one cycle to the next.  6) Episodic post-shock dust formation in the colliding winds of Eta Carinae may favor different refractory materials than those which condensed in the polar lobes.  The more massive Homunculus may have had much higher densities and lower temperatures following the 19th century Great Eruption, fostering efficient formation of silicates, Fe grains, or other species that require lower temperatures.  The transitory nature of the dust formation in colliding winds, however, may only provide sufficiently high densities for a brief time when the temperatures are high, permitting only corundum, amorphous carbon, or other species that condense at high temperatures.  The limited duration of dust formation could then preclude the formation of lower temperature condensates.  This might lead to spatial variations in dust composition and grain size around Eta Carinae, of which there are some observations signatures (e.g., Robinson et al.\\ 1987; Smith et al.\\ 1998; Chesneau et al.\\ 2005).  \\smallskip  It would obviously be interesting to spatially resolve the expanding dust that formed in recent eposides in order to constrain its geometry, as has been done for WR systems mentioned earlier, and to obtain mid-IR spectra of this newly formed dust as it expands and cools.  This may prove a difficult challenge amid the bright and spatially complex Homunculus, however.  Dust formation now may be less efficient than it was during 1981--1987, due to changes in the primary star's mass-loss.    \\smallskip\\smallskip\\smallskip\\smallskip \\noindent {\\bf ACKNOWLEDGMENTS} \\smallskip \\scriptsize  I am indebted to P.A.\\ Whitelock for providing electronic tables of the $JHKL$ photometry obtained at the SAAO from 1972 to 2004, and I thank the referee, P.\\ Williams, for helpful comments that improved the paper.  I was partially supported by NASA through grants GO-10241 and GO-10475 from the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS5-26555."
    },
    "0910/0910.3260_arXiv.txt": {
        "abstract": "Hypervelocity stars (HVSs) escaping away from the Galactic halo are dynamical products of interactions of stars with the massive black hole(s) (MBH) in the Galactic Center (GC). They are mainly B-type stars with their progenitors unknown. OB stars are also populated in the GC, with many being hosted in a clockwise-rotating young stellar (CWS) disk within half a parsec from the MBH and their formation remaining puzzles. In this paper, we demonstrate that HVSs can well memorize the injecting directions of their progenitors using both analytical arguments and numerical simulations, i.e., the ejecting direction of an HVS is almost anti-parallel to the injecting direction of its progenitor. Therefore, the spatial distribution of HVSs maps the spatial distribution of the parent population of their progenitors directly. We also find that almost all the discovered HVSs are spatially consistent with being located on two thin disk planes. The orientation of one plane is consistent with that of the (inner) CWS disk, which suggests that most of the HVSs originate from the CWS disk or a previously existed disk-like stellar structure with an orientation similar to it. The rest of HVSs may be correlated with the plane of the northern arm of the mini-spiral in the GC or the plane defined by the outer warped part of the CWS disk.  Our results not only support the GC origin of HVSs but also imply that the central disk (or the disk structure with a similar orientation) should persist or be frequently rejuvenated over the past 200 Myr, which adds a new challenge to the stellar disk formation and provides insights to the longstanding problem of gas fueling into massive black holes. ",
        "introduction": "\\label{sec:intro}  In the past several years, surveys of hypervelocity stars (HVSs) have found 16 HVSs with velocities substantially higher than the Galactic escape velocity \\citep{Brown05,Hirsch05,Edelmann05,Brown07,Brown09a}.  Most HVSs are B-type stars probably with masses $\\sim3-4\\msun$, ages $\\ga(1-2)\\times10^8$~yr, and lifetime $\\sim(1.6-3.5)\\times10^8$~yr \\citep{Brown07,Brown09a}.  Their heliocentric distances range from several ten to a hundred kpc and their spatial distribution on the sky is probably anisotropic \\citep{Brown09a,Brown09b,Abadi09}. The young nature and the spatial anisotropic distribution of these stars, as well as their hypervelocities, should be related to their origin.  A few dynamical mechanisms, involving interactions with the massive black hole (MBH) in the Galactic Center (GC), were proposed to eject stars with such hyper velocities \\citep{Hills88,YT03,Bromley06,Gualandris05,Levin06,Baumgardt06,Sesana06,Perets07,OL08,LB08}, including tidal breakup of binary star systems by the central MBH and three-body interactions of single stars with a hypothetical binary black hole (BBH) in the GC. The sources of (binary) stars injected into the vicinity of the MBH are often (implicitly) assumed to be isotropically distributed in previous studies, as the majority of the GC stars are isotropically distributed; and different ejection mechanisms may result in different spatial distributions of HVSs, depending on whether the central MBH is a single or a binary and relevant binary parameters \\citep{YT03,Bromley06,Sesana06}. Hence, the spatial distribution of HVSs was proposed to be useful in identifying the ejection mechanism responsible for the observed HVSs.\\footnote{Some other observational properties of HVSs, such as the binarity, rotational velocity and metallicity, were also proposed to be useful in identifying the ejection mechanisms of the HVSs (e.g., \\citealt{Luetal07,H07,Przybilla08,LMB08,Perets09a,Perets09b}).}  The progenitors of the discovered B-type young HVSs and their sources, however, should be distinct populations from the majority of the GC stars (which are typically old).  Observations have indeed shown various young stellar structures in the GC, which may be possible parent populations of the HVS progenitors, including a clockwise young stellar disk (CWS) and the other possible counterclockwise one within $0.5$~pc from the central MBH \\citep{LB03,Paumard06,LuJ09,Bartko09a}, young stellar clusters like the Arches and the Quintuplet systems at several ten pc away from the center, and the tidal remnants of those clusters \\citep{Portegies02}.  There also exist other organized structures in the GC, such as, the circumnuclear molecular disk, the northern arm (Narm) and the bar components of the minispiral at a few pc from the center, with which young stars may be associated \\citep{YMW00,Paumard06}. These structures may also be related to the sources of HVSs. All the structures above are not isotropically distributed, but either planar or orbiting on some specific planes around the central MBH.  The HVSs, if ejected by interactions of stars originating from these sources with the central MBH(s), are very likely to be spatially correlated. And their spatial distribution should be mapping the distribution of the parent populations of their progenitors, if the ejected HVSs can well memorize the injecting direction of their progenitors.  This paper is organized as following. In \\S~\\ref{sec:deflection}, we study how the spatial distribution of HVSs is related with the geometrical structure of the parent population of their progenitors. In \\S~\\ref{sec:Obsconst} we find that most HVSs discovered so far are probably originated from the clockwise young stellar disk (CWS) or a disk-like stellar structure with an orientation similar to that of the CWS disk in the GC. Discussion and conclusions are given in \\S~\\ref{sec:disc} and \\S~\\ref{sec:con}. ",
        "conclusions": "\\label{sec:con}  Using both analytical arguments and numerical simulations, we have demonstrated that HVSs can well memorize the injecting directions of their progenitors. In another words, the ejecting direction of an HVS is almost anti-parallel to the injecting direction of its progenitor. Therefore, the spatial distribution of HVSs should map the spatial distribution of the parent population of their progenitors directly. We also find that most of the discovered HVSs are spatially consistent with being located on two thin disk planes. The orientation of one plane is consistent with that of the (inner) CWS disk, which suggests that most of the HVSs originate from it or a disk-like stellar structure with a similar orientation to it.  The rest of HVSs may be correlated with the plane of the northern arm of the mini-spiral in the GC or the plane defined by the outer warped part of the CWS disk. Our results not only support the GC origin of HVSs but also imply that the central disk (or the disk structure with a similar orientation) should persist or be frequently rejuvenated over the past 200 Myr, which adds a new challenge to the stellar disk formation and provides insights to the longstanding problem of gas fueling into massive black holes."
    },
    "0910/0910.0922_arXiv.txt": {
        "abstract": "We present the results of $AKARI$ observations of the O-rich supernova remnant G292.0$+$1.8 using six IRC and four FIS bands covering 2.7--26.5 $\\mu$m and 50--180 $\\mu$m, respectively. The $AKARI$ images show two prominent structures; a bright equatorial ring structure along the east-west direction and an outer elliptical shell structure surrounding the remnant. The equatorial ring structure is clumpy and incomplete with its western end opened. The outer shell is almost complete and slightly squeezed along the north-south direction. The central position of the outer shell is $\\sim$ 1$'$ northwest from the embedded pulsar and coincides with the center of the equatorial ring structure. In the northen and southwestern regions, there is also faint emission with a sharp boundary beyond the bright shell structure. The equatorial ring and the elliptical shell structures were partly visible in optical and/or X-rays, but they are much more clearly revealed in our $AKARI$ images. There is no evident difference in infrared colors of the two prominent structures, which is consistent with the previous proposition that both structures are of circumstellar origin. However, we have detected faint infrared emission of a considerably high 15 to 24 $\\mu$m ratio associated with the supernova ejecta in the southeastern and northwestern areas. Our IRC spectra show that the high ratio is at least partly due to the emission lines from Ne ions in the supernova ejecta material. In addition we detect a narrow, elongated feature outside the SNR shell. We derive the physical parameters of the infrared-emitting dust grains in the shocked circumstellar medium and compare the result with model calculations of dust destruction by a SN shock. The $AKARI$ results suggest that the progenitor was at the center of the infrared circumstellar shell in red supergiant stage and that the observed asymmetry in the SN ejecta could be a result of either a dense circumstellar medium in the equatorial plane and/or an asymmetric explosion. ",
        "introduction": "Young core-collapse supernova remnants (SNRs) in the Galaxy provide an unique opportunity to study fine details of ejecta from supernova (SN) explosions as well as those of the circumstellar medium (CSM) produced over the final evolutionary stages of massive stars. An obvious group of such young core-collapse SNRs is O-rich SNRs that show optical spectra featuring strong O and Ne lines, with lines from lighter elements (e.g., He) either weak or absent. Since the ejecta from a progenitor and the swept-up CSM in the O-rich SNRs are not completely mixed together yet, observation of the O-rich SNRs is promising to study the ejecta and CSM related to the explosion of a progenitor and its late-stage evolution.  G292.0$+$1.8 (MSH11--54), together with Cassiopeia~A and Puppis~A, forms a rare group of O-rich SNRs in the Galaxy. The O-rich nature of \\snr\\ was discovered by the optical detection of fast-moving O-rich and Ne-rich ejecta \\citep{gos79, mur79}. Recent wide-field observations of the optical \\othree\\ line in \\snr\\ have revealed that the ejecta velocities range from $-$1400~\\kms\\ to $+$1700~\\kms\\ and also that distant ejecta knots are located primarily along the north-south direction \\citep{gha05, win06}. The dynamical center of G292.0$+$1.8, based on the kinematics of the \\othree\\ line emission from the ejecta, is roughly coincident with the geometrical center of the radio emission, but shows an apparent offset from its pulsar position discovered at the southeast \\citep[][see below]{win08}. \\cite{gha09} reported that the mid-infrared (MIR) spectrum of the ejecta is dominated by ionic lines and a broad bump around 17~$\\mu$m using the $Spitzer$ Space Telescope ($Spitzer$). They proposed that the latter is produced either by Polycyclic Aromatic Hydrocarbons (PAHs) along the line of sight or newly-formed dust within the ejecta. Besides there is also an equatorial MIR continuum component from the dust associated with the shocked CSM partly overlapping with the ejecta emission. A previous infrared (IR) study based on the data obtained with Infrared Astronomical Satellite ($IRAS$), on the other hand, identified enhanced far-IR (FIR) emission in the southwestern part of the SNR, suggesting that the source has encountered adjacent clouds \\citep{bra86, par07}.  In radio, \\snr\\ mainly consists of bright emission from a central pulsar wind nebula (PWN) of $\\sim$~4\\arcmin\\ (in diameter) and relatively fainter outer plateau emission of $\\sim$~9\\arcmin\\ \\citep{loc77, bra86, gae03}. The plateau emission declines sharply outward, and there is no apparent surrounding radio shell structure. The pulsar PSR J1124--5916, which is located $\\sim$46\\arcsec\\ from the dynamical center of the SNR in the southeast direction, was discovered by radio timing \\citep{cam02}, confirming the PWN nature of the central emission of the SNR. The pulsar and PWN were also identified in X-ray emission \\citep{hug03}, where the pulsar has a nearby jet and a torus of $\\sim$~5\\arcsec\\ scale \\citep{par07}. The jet axis is slightly tilted to the northeast direction \\citep{par07}. There are two clear differences between the soft and hard X-ray emission of \\snr. First, while the former is produced by shocked gas of normal composition distributed as equatorial filaments, indicating it is likely from CSM, the latter is dominated by bright metallic lines from the ejecta \\citep{par02, par04}. Next, the X-ray emission is harder in the northwestern part than the southeastern part, implying density variation or asymmetric SN explosion \\citep{par07}. The distance to \\snr, on the other hand, was estimated to be 6.2 kpc from H~I absorption observations \\citep{gae03}.  In this paper we present an extensive IR study of G292.0$+$1.8 using multi-band imaging and spectroscopic data obtained with the $AKARI$ space telescope \\citep{mur07}. We describe the details of our $AKARI$ observations and present the results of the observations in \\S2 and 3, respectively. In \\S4 we discuss IR emission in \\snr\\ from circumstellar dust with a comparison to model calculations of dust destruction by a SN shock. We also discuss the IR emission from the ejecta and we suggest that there is a negligible amount of dust associated with the ejecta. Then, we make a comparison to the results of $Spitzer$ obserations. We finally discuss the SN explosion in \\snr\\ based on our $AKARI$ observations followed by our conclusions in \\S5. ",
        "conclusions": "We have presented the NIR to FIR imaging and MIR spectroscopic observations of the O-rich SNR G292.0$+$1.8 using the IRC and FIS instruments aboard $AKARI$ satellite. The almost continuous multiband imaging capability of $AKARI$ covering wide IR wavelengths together its wide field of view enabled us to clearly see the distinct IR emission from the entire SNR and to derive its IR charactersitics. We derive the physical parameters of IR-emitting dust grains in the swept-up circumstellar medium and compare the result with the model calculations of dust destruction by a SN shock. The overall shape of the observed SED can be explained by a simple model using characteristic SNR parameters with a bit lower initial dust-to-gas ratio. At 11 $\\mu$m, the model flux is significantly smaller, which may indicate the importance of stochastic heating. We have not detected any signficant amount of freshly-formed dust associated with the SN ejecta.  The AKARI results in this paper give new insights into the explosion dynamics of G292.0$+$1.8. We have discovered an almost symmetric IR shell probably produced by the circumstellar wind from the progenitor star in the RSG phase. Its center is significantly offset from the previously suggested explosion centers. We consider that either OES represents the circumstellar shell pressure-confined by external medium or the SN exploded close to the center of the  OES. In the latter case, the pulsar in G292.0$+$1.8 may be traveling at a speed of $\\sim 1,000$~\\kms. At the same time, the ejecta distribution is unveiled by their high 15 to 24 $\\mu$m ratio. The ejecta are mainly distributed along the northwest-southeast direction. This symmetric pre-supernova structure and asymmetric ejecta distribution appear to be rather common in the remnants of SN IIL/b, which suffer strong mass-loss like Cas A \\citep{hin04, smi08}. There is also a Narrow tail outside the SNR shell, which might be similar to the feature observed in Cas~A. A detailed study is necessary to understand the nature of this feature.  We suggest that multi-band IR imaging observations are powerful tools to explore both the ejecta and CSM emission in young core-collapse SNRs. Especially the 15 and 24 $\\mu$m images are useful to reveal the detailed structure of IR features, which leads to better understanding of environments of the progenitor and the SN explosion.   This work is based on observations with $AKARI$, a JAXA project with the participation of ESA. We wish to thank all the members of the $AKARI$ project. We also thank B. Gaensler for providing the ATCA 20 cm image and S. Park for providing the $Chandra$ X-ray image. This work was supported by the Korea Research Foundation Grant funded by the Korean Government (KRF-2008-357-C00052) and the Korea Science and Engineering Foundation (R01-2007-000-20336-0). This work was also supported in part by a Grant-in-Aid for Scientific Research for the Japan Society of Promotion of Science (18204014). T.N.\\ has been supported in part by World Premier International Research Center Initiative (WPI Initiative), MEXT, Japan, and by the Grant-in-Aid for Scientific Research of the Japan Society for the Promotion of Science (19740094).  {\\it Facility:} \\facility{Akari}"
    },
    "0910/0910.3676_arXiv.txt": {
        "abstract": "Recently, a new class of radio transients in the 5-GHz band and with durations of the order of hours to days, lacking any visible-light counterparts, was detected by Bower and collaborators. We present new deep near-Infrared (IR) observations of the field containing these transients, and find no counterparts down to a limiting magnitude of $K=20.4$\\,mag. We argue that the bright ($>1$\\,Jy) radio transients recently reported by Kida et al. are consistent with being additional examples of the Bower et al. transients. We refer to these groups of events as ``long-duration radio transients''. The main characteristics of this population are: time scales longer than 30\\,minute but shorter than several days; very large rate, $\\sim10^{3}$\\,deg$^{-2}$\\,yr$^{-1}$; progenitors sky surface density of $>60$\\,deg$^{-2}$ (at $95\\%$ confidence) at Galactic latitude $\\sim40^{\\circ}$; 1.4--5\\,GHz spectral slopes, $f_{\\nu}\\propto \\nu^{\\alpha}$, with $\\alpha\\gtorder0$; and most notably the lack of any X-ray, visible-light, near-IR, and radio counterparts in quiescence. We discuss putative known astrophysical objects that may be related to these transients and rule out an association with many types of objects including supernovae, gamma-ray bursts, quasars, pulsars, and M-dwarf flare stars. Galactic brown-dwarfs or some sort of exotic explosions in the intergalactic medium remain plausible (though speculative) options. We argue that an attractive progenitor candidate for these radio transients is the class of Galactic isolated old neutron stars (NS). We confront this hypothesis with Monte-Carlo simulations of the space distribution of old NSs, and find satisfactory agreement for the large areal density. Furthermore, the lack of quiescent counterparts is explained quite naturally. In this framework we find: the mean distance to events in the Bower et al. sample is of order kpc; the typical distance to the Kida et al. transients are constrained to be between 30\\,pc and 900\\,pc (at the 95\\% confidence level); these events should repeat with a time scale of order several months; and sub-mJy level bursts should exhibit Galactic latitude dependence. We discuss two possible mechanisms giving rise to the observed radio emission: incoherent synchrotron emission and coherent emission. We speculate that if the latter is correct, the long duration radio transients are sputtering ancient pulsars or magnetars and will exhibit pulsed emission. ",
        "introduction": "\\label{sec:Introduction}  Large field-of-view radio telescope facilities such as the Parks multi-beam facility, the Arecibo multi-beam instrument, the Allen Telescope Array (DeBoer et al. 2004), and the Low Frequency Array (Falcke et al. 2007) have reinvigorated the radio-frequency time domain frontier. First signs of this ``revolution'' are indicated by the discoveries of new classes of radio transients. Examples include: several Galactic center sources (e.g., Hyman  et al. 2005; 2009); Rotating Radio Anomalous Transients (RRATs; McLaughlin et al. 2006) which represent a previously unknown class of radio pulsars, probably twice as abundant as their ``normal'' cousins; and the powerful ($\\sim30$~Jy) radio ``Sparker'', with a time scale of several milliseconds (Lorimer et al. 2007; Kulkarni et al. 2009).   Here, we focus on yet another emerging class of mysterious radio transients. In a novel approach, Bower et al. (2007) re-analyzed 944 epochs of Very Large Array\\footnote{The Very Large Array is operated by the National Radio Astronomy Observatory, a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.} (VLA) observations, taken about once per week for twenty two years, of a single calibration field. These authors discovered a total of ten transients, nine in the 5-GHz band and one in the 8-GHz band. These transients can be divided into two groups: ``single-epoch'' and ``multi-epoch'' transients. The eight single-epoch transients, as can be gathered by their names, were detected in only one epoch. Given a single epoch detection, one can only constrain the duration of the transient by the epochs preceding and succeeding the time at which the transient was detected (approximately one week). The lower limit could be as small as the typical integration time (about 20\\,minute). The two multi-epoch transients were detected after averaging over two months of data. Thus, the duration of these two events can be taken to be about two months.  Bower et al. (2007) split the data for each epoch, consisting of 20 minutes, into five segments and looked for variability on four-minute time scale. They did~not find any evidence for variability on these time scales. However, the total S/N of these detections was $\\ltorder7$, and therefore the limits on variability within these 20-minute windows are appropriately weak. More importantly, the authors find the circular polarization is less than $\\sim30\\%$.  Separately, Kuniyoshi et al. (2006), Niinuma et al. (2007), and Kida et al. (2008) reported on a search for radio transients using an East-West interferometer of the Nasu Pulsar Observatory (located in Tochigi Prefecture, Japan) of Waseda University. The program consists of daily drift scanning of the sky towards the local zenith. These authors reported six bright radio transients (and several other were mentioned but without details), with flux density above 1\\,Jy in the 1.4-GHz band. Five were single epoch transients while one was detected on two successive days with flux densities of 1.7 and 3.2\\,Jy, respectively (Niinuma et al. 2007). In each epoch the transients were detected for about 4\\,minutes, which is the drift scanning time, and did~not exhibit any significant variation within the observation. Unfortunately, these events are not well localized and have positional uncertainties of the order of $0.4^{\\circ}$ in declination, and $0.04^{\\circ}$ in right ascension. Kida et al. (2008) stated that the 2-$\\sigma$ upper limit for the rate of these transients is $0.0049$\\,deg$^{-2}$\\,yr$^{-1}$. Arguably, these events also have time scales somewhere in the range of a few minutes to a few days. Therefore, later we consider the framework in which both the VLA and the Nasu events have a common origin.  Another relevant survey was conducted by Levinson et al. (2002) who compared the NRAO VLA Sky Survey (NVSS; Condon et al. 1998) and the ``Faint Images of the Radio Sky at Twenty centimeters'' survey (FIRST; Becker et al. 1995; White et al. 1997). Both these surveys were undertaken in the 1.4-GHz band and their 5-$\\sigma$ limits are 3.5 and 1\\,mJy, respectively. Levinson et al. (2002) identified nine radio transient candidates with flux densities greater than 6\\,mJy. Followup observations of these radio transients (Gal-Yam et al. 2006) showed that seven were spurious and the remaining two were plausible transients\\footnote{The term ``transient'' is used here in the sense that we do~not detected emission in quiescence.}: an optically extincted SN in NGC~4216 from which the radio emission lasted for several years; and VLA~J172059.9$+$385229 (discussed in \\S\\ref{sec:VLA1720}).     Bower et al. (2007) obtained deep visible light images of their transients. They found that the multi-epoch event RT\\footnote{Here RT stands for radio transient and the succeeding eight digits are yyyymmdd where yyyy is the year, mm is the month and dd is the day.}\\,19870422 was 1\\farcs5 from a $z=0.249$, $R=20.2$\\,mag galaxy. The peak luminosity of RT\\,19870422, assuming that the transient is related to this galaxy, is consistent (to an order of a magnitude) with an energetic supernova (SN) similar to SN\\,1998bw (GRB\\,980425; Kulkarni et al. 1998) and SN\\,2006aj (Soderberg et al. 2006). We note that the rate of these events is marginally consistent with the rate of low-luminosity GRBs derived by Soderberg et al. (2006). In contrast, the other multi-epoch transient RT\\,20010331 has no optical counterpart to a limiting magnitude of $g\\approx27.6$\\,mag, $R\\approx26.5$\\,mag, and $K\\approx19.2$\\,mag (this paper) to within $5''$ of the radio source. The great offset between a putative host galaxy and the radio transient make this an unusual source. We note that Cenko et al. (2008) presented an example of a GRB in a galaxy halo environment. However, only about $1\\%$ of all GRBs occures in such environments.    The single-epoch radio transient RT\\,19840613 falls within the optical boundary of a $z=0.040$ spiral galaxy, but clearly lying outside the nucleus of the galaxy. On the basis of the radio luminosity (assuming association with the galaxy) and the nature of the putative host galaxy this transient is consistent with an origin similar to that of RT\\,19870422 (i.e., a low luminosity GRB).  The remaining seven single-epoch transients (one detected in 8\\,GHz and six at 5\\,GHz) do not have astrometrically coincident optical, near-IR, or radio counterparts and have a point source appearance (see Table~\\ref{tab:FastTran}). This is a major clue in that a large fraction of gamma ray bursts (GRBs) and most SNe have detectable optical host galaxies at $R\\sim26$\\,mag level (e.g., Ovaldsen et al. 2007).    In order to separate the events discussed above from Sparkers (Lorimer et al. 2007; Kulkarni et al. 2009), which have very short time scales, we refer to these events as ``long-duration radio transients''.   In Table~\\ref{tab:FastTran} we summarize the observational properties of all the long-duration radio transients. We define this class as events with no optical identification and durations between hours to days. This group include the seven known sources from Bower et al. (2007) that are not associated with any optical counterpart and have time scales shorter than about one week; RT\\,19870422 (see above); and the six bright transients reported by Kida et al. (2008).   The structure of this paper is as follows. In \\S\\ref{sec:VLA1720} we re-examine the case of VLA\\,J172059.9$+$385229 (Levinson et al. 2002) and show it is a spurious event. In \\S\\ref{sec:Obs} we present new near-IR observations of the Bower et al. field. In \\S\\ref{sec:Prop} we review the basic properties (areal density, annual rate) of the long-duration radio transients. Next, in \\S\\ref{sec:NotProgenitor} we use the observational clues to refute several plausible explanations regarding the nature of the long-duration radio transients. In \\S\\ref{sec:IONS} we argue that the most attractive explanation is that these radio transients are associated with Galactic isolated old Neutron Stars (NS). Finally, we discuss and summarize the results in \\S\\ref{sec:Disc}.   \\begin{deluxetable*}{llllllllllllll} \\tablecolumns{14} \\tablewidth{0pt} \\tablecaption{List of candidate long-duration radio transients} \\tablehead{ \\multicolumn{9}{c}{Transient} & \\multicolumn{4}{c}{Lim. mag.} & \\colhead{} \\\\ \\colhead{Epoch} & \\colhead{R.A.} & \\colhead{Dec.} & \\colhead{Band} & \\colhead{$S$} & \\colhead{$\\delta{t}$\\tablenotemark{a}} & \\colhead{$S_{next}$\\tablenotemark{b}} & \\colhead{$S_{deep}$\\tablenotemark{c}} & \\colhead{$S/S_{deep}$} & \\colhead{X\\tablenotemark{d}} & \\colhead{$g$} & \\colhead{$R$} & \\colhead{$K$} & \\colhead{Ref.} \\\\ \\colhead{}        & \\colhead{J2000}        & \\colhead{J2000}        & \\colhead{GHz}     & \\colhead{$\\mu$Jy} & \\colhead{day}     & \\colhead{$\\mu$Jy} & \\colhead{$\\mu$Jy} & \\colhead{}     & \\colhead{cts}  & \\colhead{mag}  & \\colhead{mag}  & \\colhead{mag}  & \\colhead{} } \\startdata 1984 05 02 & 15 02 24.61& $+$78 16 10.1 & 5.0 & $448\\pm74$   & 7  & $-10\\pm68$   & $<8$  & 56  & 0.08 & 27.6 & 26.5 & 20.4 & 1\\\\ 1986 01 15 & 15 02 26.40& $+$78 17 32.4 & 5.0 & $370\\pm67$   & 7  & $199\\pm121$  & $<8$  & 46  & 0.07 & 27.6 & 26.5 & 20.4 & 1\\\\ 1986 01 22 & 15 00 50.15& $+$78 15 39.4 & 5.0 & $1586\\pm248$ & 7  & $-59\\pm164$  & $<15$ & 106 & 0.08 & 27.6 & 26.5 & 20.2 & 1\\\\ 1992 08 26 & 15 02 59.89& $+$78 16 10.8 & 5.0 & $642\\pm101$  & 56 & $37\\pm83$    & $<9$  & 71  & 0.08 & 27.6 & 26.5 & 19.6 & 1\\\\ 1997 05 28\\tablenotemark{e} & 15 00 23.55& $+$78 13 01.4 & 5.0 & $1731\\pm232$ & 7  & $90\\pm206$   & $<36$ & 48  & 0.07 & 27.6 & 26.5 & 20.0 & 1\\\\ 1999 05 04 & 14 59 46.42& $+$78 20 29.0 & 5.0 & $7042\\pm963$ & 21 & $-313\\pm1020$& $<117$& 60  & 0.05 & 27.6 & 26.5 & 19.2 & 1\\\\ \\hline 1997 02 05 & 15 01 29.35& $+$78 19 49.2 & 8.4 & $2234\\pm288$ & 5  & $857\\pm323$  & $<646$& 3.5 & 0.08 & 27.6 & 26.5 & 20.4 & 1\\\\ \\hline 2005 01 10 & 04 45 17   & $+$41 30      & 1.4 & $1.8\\times10^{6}$ & 1 & $\\ltorder3\\times10^{5}$ & $\\ltorder9\\times10^{3}$ & 200 & 0.02 & & & & 2 \\\\ 2005 03 27 & 06 45 15   & $+$32 00      & 1.4 & $1.2\\times10^{6}$ & 1 & $\\ltorder3\\times10^{5}$ & $\\ltorder6\\times10^{3}$ & 200 & 0.02 & & & & 2 \\\\ 2005 03 04 & 10 39 43   & $+$32 00      & 1.4 & $1.7\\times10^{6}$ & 1 & $\\ltorder3\\times10^{5}$ & $\\ltorder7\\times10^{3}$ & 240 & 0.04 & & & & 2 \\\\ 2005 01 02 & 10 43 06   & $+$41 00      & 1.4 & $1.7\\times10^{6}$ & 1 & $\\ltorder3\\times10^{5}$ & $\\ltorder1\\times10^{3}$ & 1700& 0.02 & & & & 2 \\\\ 2005 02 13\\tablenotemark{f} & 14 43 22   & $+$34 39      & 1.4 & $3.2\\times10^{6}$ & 1 & $\\ltorder3\\times10^{5}$ & $\\ltorder2\\times10^{3}$ & 1600 & 0.02 & & & & 2 \\\\ 2004 03 20 & 17 37 17   & $+$38 08      & 1.4 & $1.0\\times10^{6}$ & 1 & $\\ltorder3\\times10^{5}$ & $\\ltorder6\\times10^{3}$ & 170 & 0.02 & & & & 2 \\\\ \\hline \\hline 2001 10 31\\tablenotemark{g} & 15 03 46.18& $+$78 15 41.7 & 4.0 & $697\\pm94$   & 59 & $85\\pm85$    & $<37$ & 19  & 0.08 & 27.6 & 26.5 & 19.2 & 1 \\enddata \\tablenotetext{a}{Time to next observation.} \\tablenotetext{b}{Flux limit in the next observation.} \\tablenotetext{c}{ Specific flux limit on radio emission at quiescence. For the Bower et al. transients this limit is obtained from the non detection in the combined image of the Bower et al. field. For the Kida et al. (2008) transients we list the flux of the brightest FIRST or NVSS radio source (or detection limit if no source) in the transient positional error region.} \\tablenotetext{d}{ROSAT 3-$\\sigma$ upper limit (in count\\,s$^{-1}$) in the 0.12-2.48\\,keV band. For the Kida et al. (2008) transients we list the count rate of the brightest ROSAT source (or detection limit if no source) in the transient positional error region. We obtained these limits by calculating the 3-$\\sigma$ noise level due to the background in the ROSAT X-ray images at the location of each transient.} \\tablenotetext{e}{A galaxy with $R=19.6$\\,mag and $z=0.245$, $5''$ away, probably due to chance coincidence.} \\tablenotetext{f}{Detected on two epochs, separated by one day, with fluxes of 1.7\\,Jy and 3.2\\,Jy on the first and second epochs, respectively.} \\tablenotetext{g}{RT\\,20011031 had a time scale of two~months and is listed here for completeness.} \\tablecomments{A list of candidate long-duration radio transients and their properties. Transients detected in different frequencies or instruments are separated by horizontal lines. The first (second) block lists the six (one) single-epoch transients detected by Bower et al. (2007) at 5\\,GHz (8\\,GHz) with no optical counterpart. The third block lists the Kida et al. (2008) events, and the fourth block lists the two-months 5\\,GHz event detected by Bower et al. (2007). This last event is shown here for completeness. References: (1) Bower et al. (2007); (2) Kida et al. (2008). We note that the Kida et al. (2008) transients have positional uncertainties of order $0.4^{\\circ}$ in declination, and $0.04^{\\circ}$ in right ascension.} \\label{tab:FastTran} \\end{deluxetable*} ",
        "conclusions": "\\label{sec:Disc}  We review several recent discoveries of radio transients with durations between minutes to days (Kuniyoshi et al. 2006; Bower et al. 2007; Niinuma et al. 2007; Kida et al. 2008). We suggest that these radio transients may be generated by a single class of progenitors. The main characteristics of these ``long duration radio transients'' are: (i) a very high occurrence rate of about $\\sim10^{3}$\\,deg$^{-2}$\\,yr$^{-1}$ in the 5-GHz band; (ii) common at intermediate Galactic latitude, with progenitors sky surface density of $\\gtorder60$\\,deg$^{-2}$ at Galactic latitude, $b\\sim40^{\\circ}$; (iii) lacking any X-ray ($\\gtorder10^{-13}$\\,erg\\,cm$^{-2}$\\,s$^{-1}$), visible light ($g<27.6$, $R<26.5$\\,mag), near-IR ($K<20.4$\\,mag) and radio ($S_{5\\,{\\rm GHz}}\\gtorder10\\,\\mu$Jy) counterparts; and (iv) more abundant in the 5-GHz band as compared to that in the 1.4-GHz band; From the rates in the two bands we infer the spectral index between the two bands is $\\alpha \\gtorder 0$ where the flux density, $f_{\\nu} \\propto \\nu^{\\alpha}$.   These events are most probably not associated with the usual culprits like GRBs, SGRs, AGNs, SNe, flare stars, pulsars and interacting binaries (\\S\\ref{sec:NotProgenitor}). We find that several other hypothesis including Galactic isolated old NS; brown dwarfs and some sort of a new kind of explosions cannot be ruled out. Among these, we find that the association with isolated old NSs is especially attractive. We explore this hypothesis in details and show that it is consistent with the current observations. In the framework of Galactic isolated old NSs we show that: (i) the typical distance to the Bower et al. sample (mJy events) is between 1\\,kpc and $\\sim5$\\,kpc; (ii) the typical distance to the Kida et al. Jy-level events is less than about 0.9\\,kpc and more than 30\\,pc at the $95\\%$ CL; and (iii) they will have a burst repetition time scale of about several months and duty cycle of $\\sim7\\times 10^{-3}$.  A possible association with isolated old NS is exciting. If correct, this may prove to be the most practical way, so far, to find old NSs in our Galaxy and explore their demographics. Specifically, the demography of old NSs constitute an excellent probe of the star formation history and metal enrichment of the Galaxy, and the gravitational potential of the Galaxy.    Our analysis naturally suggests several tests that can be used to rule out this hypothesis. First, if indeed long-duration radio transients are associated with isolated old NSs, then we expect them to be distributed inhomogeneously on the celestial sphere. Specifically, this hypothesis predicts that faint sub-mJy long-duration radio transients will be at least a few times more common in the Galactic center than in high Galactic latitude. The exact ratio, however, depends on the distance scale to these events. This is illustrated in Figure~\\ref{fig:NS_SkyDistribution} which shows the expected approximate distribution of isolated old NSs on the celestial sphere, for the entire NS population and NS which are at distance smaller than 1\\,kpc from the Sun. Finally, we note that if the radio emission from such sources is pulsating, then pulsars searches conducted at 5\\,GHz will find these objects. The fact that such ``pulsars'' were not found in existing surveys may be due to the fact that the majority of pulsars searches are carried on in low frequencies in which the rate of these transients is low."
    },
    "0910/0910.4260_arXiv.txt": {
        "abstract": "{ We present an explanation of the 4th branch of the Z-track based on analysis of high-quality {\\it RXTE} data on the source GX\\th 5-1. Spectral analysis shows that the physical evolution on the 4th track is a continuation of the flaring branch which we previously proposed consists of unstable nuclear burning of the accretion flow on the neutron star. In flaring there is a huge increase of the neutron star emission from a volume that increases to a radius of 21 km. The 4th branch is shown to consist of flaring under conditions that the mass accretion rate and thus the total source luminosity is falling. We detect strong emission on the flaring and 4th branches at energies between 7.8 - $\\sim$9.4 keV inconsistent with origin as Fe K emission, which we suggest is the radiative recombination continua (RRC) of iron Fe XXVI at 9.28 keV and of lower states. Evolution of the emission takes place, the energy falling but the flux increasing strongly, consistent with production in the large volume of unstable nuclear burning around the neutron star which eventually cools. ",
        "introduction": "The Z-track sources form the brightest group of Low Mass X-ray Binary (LMXB) sources containing a neutron star, including: GX\\th 340+0, GX\\th 5-1, Cyg\\th X-2, Sco\\th X-1, GX\\th 17+2 and GX\\th 349+2 (Hasinger \\& van der Klis 1989). All of these persistently radiate at or above the Eddington limit for a neutron star. The sources trace a Z-shaped pattern in X-ray hardness versus intensity consisting of the horizontal branch (HB), the normal branch (NB) and the flaring branch (FB). The sources are divided into the Cyg-like sources: GX\\th 340+0, GX\\th 5-1 and \\hbox{Cyg\\th X-2} with long horizontal branches and the Sco-like sources: Sco\\th X-1, GX\\th 17+2 and GX\\th 349+2 with short horizontal branches but stronger flaring branches. This strong hardness evolution along the Z-track is clearly due to major physical changes taking place within the sources, the nature of which has not been understood. \\begin{figure*}                                                           % \\includegraphics[width=44mm,height=135mm,angle=270]{balucinska-church_2009_02_fig01.ps} % \\caption{Lightcurve of the observation in the energy band 1.9 - 18.5 keV. The data consist of 4 sub-observations so that, to improve clarity, the gaps between the sub-observations have been compressed.} \\end{figure*} However, based on analysis of {\\it Rossi-XTE} data on the sources GX\\th 340+0, GX\\th 5-1 and Cygnus\\th X-2, we have proposed a model of the Cyg-like sources (Church et al. 2006; Jackson et al. 2008, Ba\\l uci\\'nska-Church et al. 2009). In this model, the increase in X-ray intensity moving on the NB from the soft apex to the hard apex is due to an increase of mass accretion rate. There is a resultant heating of the neutron star and a contraction in the emitting area so that the flux emitted per unit area of the neutron star surface becomes super-Eddington leading to disruption of the inner disk and the launching of jets detected as radio emission on this part of the Z-track. On the FB, spectral fitting showed little change of $\\dot M$ but the large increase in the luminosity of the neutron star indicated unstable nuclear burning on the neutron star. Moreover there was good agreement at the onset of flaring (the soft apex) between the mass accretion rate per unit area of the neutron star $\\dot m$ and the theoretical value below which burning of the accretion flow on the neutron star becomes unstable.  In addition, the Cyg-like sources sometimes display extra unusual tracks which similarly have not been understood. In Fig. 2 we show the additional 4th branch in GX\\th 5-1 in the observation analysed in the present work. It can be seen that this begins at the end of the FB and moves in the direction of decreasing hardness and decreasing intensity. It has been referred to as the ``dipping track'' because of the decreasing intensity (e.g. Kuulkers \\& van der Klis 1996) by analogy with the dipping class of LMXB in which absorption in the bulge in the outer disk causes orbital-related decreases of intensity. This absorption may cause a hardening of the spectrum, but energy-independent dipping (Church \\& Ba\\l uci\\'nska-Church 1993) and also softening (Church \\& Ba\\l uci\\'nska-Church 1995) in dips is also observed. This range of behaviour may be explained in terms of two continuum emission components within the sources (Church \\& Ba\\l uci\\'nska-Church 1995). The apparent dipping in a Z-track source led to the suggestion that the Cyg-like sources have high inclination and that this explains the difference between the Cyg-like and Sco-like sub-groups. It should be noted, however, that we would not expect the 4th track to be associated with flaring but to take place on any part of the Z-track.  In the present work, we carry out an investigation of the 4th track in GX\\th 5-1. ",
        "conclusions": "Our analysis of GX\\th 5-1 shows that the 4th branch is simply a continuation of the flaring branch in which the behaviour of the neutron star blackbody temperature $kT_{\\rm BB}$ and radius $R_{\\rm BB}$ continues on the 4th branch as on the flaring branch, as shown in Fig. 4. On the 4th branch, $L_{\\rm ADC}$ decreases showing that $\\dot M$ decreases, so that the increase of $L_{\\rm BB}$ can only be due to nuclear burning on the surface of the neutron star. We previously showed that in the three Cygnus\\th X-2 like sources (e.g. Church et al. 2006), at the soft apex, the value of $\\dot m$, the mass accretion rate per unit area of the neutron star, agreed well with the critical value below which nuclear burning is unstable (e.g. Bildsten 1998). Thus as the source moves along the NB to the soft apex, unstable burning starts and the FB is formed. The presence of a 4th branch shows that $\\dot M$ continues to fall and the 4th branch is a combination of unstable burning and this.   The observation of strong emission at energies between 7.8 and $\\sim$9.4 keV is of great interest. We suggest that this is RRC emission of ionization states Fe XXVI and lower levels. Recombination of H-like Fe to He-like Fe produces RRC at 9.28 keV consistent with the peak emission energy observed. The emission is very strong which may be due to the large volume of emitting plasma  around the neutron star suggested by the large values of blackbody radius of up to 21 km. The decrease of line energy towards the end of the 4th track indicates cooling as does the fall of blackbody temperature at the end of the 4th track.  The detection of Fe RRC is very rare in Galactic sources. However, a broad line at $\\sim$10 keV was previously seen in GX\\th 5-1 on the flaring and 4th branch by Asai et al. (1994), probably the same feature as we have detected. They considered possible origin of the line as Fe K$\\alpha$, shifted to higher energies, either by a Doppler shift in the disk, or by Compton scattering or by blue-shifting in a jet, but none of these could provide a satisfactory explanation.  A strong, broad feature at 9.28 keV has previously been seen in gamma ray bursts. In GRB\\th 970828, a broad emission feature at 4.8 keV corrected for the well-determined redshift of the host galaxy of z = 0.9578 corresponds closely to 9.28 keV (Yoshida et al. 2001). $kT_{\\rm e}$ in this case was 0.8 keV producing a very broad feature. Weth et al. (2000) considered conditions under which the emission at 9.28 keV can be strong. In Galactic sources, narrow RRC emission of O VIII and O VII was seen in the dipping source XBT\\th 0748-676 (Cottam et al. (2001) and in planetary nebulae (Nordon et al. 2009) in which $kT_{\\rm e}$ is only a few eV.  In our case, the peak energy on the 4th branch corresponds to Fe XXVI RRC, i.e. H-like iron and it appears that this emission is associated with the fireball around the neutron star resulting from uncontrolled nuclear burning. Further detailed study is clearly required."
    },
    "0910/0910.0991_arXiv.txt": {
        "abstract": "The Broad Emission Lines (BELs) in spectra of type 1 Active Galactic Nuclei (AGN) can be very complex, indicating a complex Broad Line Region (BLR) geometry. According to the standard unification model  one can expect an accretion disk around a supermassive black hole in all AGN. Therefore, a disk geometry is expected in the BLR. However, a small fraction of BELs show double-peaked profiles which indicate the disk geometry. Here, we discuss a two-component model, assuming an emission from the accretion disk and one additional emission from surrounding region. We compared the modeled BELs with observed ones (mostly broad H$\\alpha$ and H$\\beta$ profiles) finding that the model can well describe single-peaked and double-peaked observed broad line profiles. ",
        "introduction": "The detailed structure of the innermost part of AGN is still an open question. It is widely accepted that the central engine consists of a super massive black hole fueled by an accretion disk. The origin of the broad emission lines, observed in most AGN, is explained by photoionization of dense gas surrounding the nuclei, but the structure and dynamics of the BLR remains unclear. Using Broad Emission Lines (BEL) shape investigations of the BLR, one can constrain the geometry of the material in the BLR \\citep[see][]{Pop06}.  BELs of AGN show high diversity in their shapes and widths. In some rare cases, where the accretion disk clearly contributes to the BELs, we can investigate the disk properties using a broad, double-peaked, low-ionization lines \\citep{Che1,Che2,Era, Era1,Era09}. The rotation of the material in the disk results in one blue-shifted and one redshifted peak, while the gravitational redshift produces a displacement of the center of the line and a distortion of the line profile.  Moreover, a number of BELs show shoulders-like profile in the wings of the line, which could be a fingerprint of the accretion disk emission \\citep{Pop03,Era1, Str03}. Beside  emission of the  disk (or disk-like region) or emission from spiral shock waves within a disk \\citep{Che1,Che2}, the following geometries may cause substructures in line profiles: i) emission from the oppositely-directed sides of a bipolar outflow \\citep{Zhen90,Zhen91}; ii) emission from a spherical system of clouds in randomly inclined Keplerian orbits illuminated anisotropically from the center \\citep{GW96}; and iii) emission of the binary black hole system \\citep{Gas83,Gas96};  To explain the complex morphology of the observed BEL shapes, different geometrical models have been discussed \\citep[see in more details][]{Sul00}. In some cases the BEL profiles can be explained only if two or more kinematically different emission regions are considered   \\citep[see e.g.][]{Marziani93,Rom96, Pop01, Pop02, Pop03, Pop04, Bon06, Ilic06, Col06, Hu08}. In particular, the existence of a Very Broad Line Region (VLBR) with random velocities at ~5000-6000 km/s within an Intermediate Line Region (ILR) has also been considered to explain the observed BEL profiles \\citep{CB96,Sul00,Hu08}.   Nevertheless, the majority of AGN with BELs have only single peaked lines, but this does not necessarily indicate that the contribution of the disk emission to the BELs profiles is negligible.   In this paper we study the possibility that the flux of the hidden disk emission could be present in single peaked  emission line profiles of AGN. With this goal, we applied a simple, two-component model, which assumes that observed BELs can be represented by composition of the emission of the accretion disk and surrounding non-disk isotropic region. For the accretion disk emission, we have used the model proposed by \\citet{Che1}, which assumes optically thick and geometrically thin disk, while for the additional emitting region which surrounds the disk, we assumed isotropic velocity distribution of emitting clouds, which imply that the emission line profiles generated by this region can be described by a Gaussian function. In Section 2 of this paper we describe samples of galaxies where we analyzed the possible accretion disk contribution to the total flux of their broad emission lines. We review the methods of analyzes in Sections 3 and in Section 4 we give results. Finally, in Section 6 we outline our conclusions. ",
        "conclusions": "In this paper we outline our recent investigation about the presence of the disk emission in single peaked BELs.  We {carried out the} following analyzes: i) fitting of single peaked BELs with Gaussians and with two component model, and ii) comparing line parameters of a sample with simulated profiles.  From our investigations we can conclude that there could be a high probability that the disk emission flux might be present in the single peaked emission line profiles. In principle, the contribution of the disk emission to the total flux of the single peaked BELs is in most cases smaller then 50\\%. Also, we found that the disk inclinations were mainly smaller then  $i<25^\\circ$. Such small inclinations and large discrepancies from simulated profiles for higher inclinations should be discussed in the contest of unified model, i.e. possible disk orientation to the torus or partial obscuration by the torus."
    },
    "0910/0910.0511_arXiv.txt": {
        "abstract": "We present the data processing and analysis techniques we are using to determine structural and photometric properties of galaxies in our Gemini/HST Galaxy Cluster Project sample.  The goal of this study is to understand cluster galaxy evolution in terms of scaling relations and structural properties of cluster galaxies at redshifts $0.15 <$ z $< 1.0$.   To derive parameters such as total magnitude, half-light radius, effective surface brightness, and Sersic n, we fit r$^{1/4}$ law and Sersic function 2-D surface brightness profiles to each of the galaxies in our sample.  Using simulated galaxies, we test how the assumed profile affects the derived parameters and how the uncertainties affect our Fundamental Plane results.  We find that while fitting galaxies which have Sersic index n $< 4$ with r$^{1/4}$ law profiles systematically overestimates the galaxy radius and flux, the combination of profile parameters that enter the Fundamental Plane has uncertainties that are small. Average systematic offsets and associated random uncertainties in magnitude and $\\log r_e$ for n $> 2$ galaxies fitted with r$^{1/4}$ law profiles are $-0.1\\pm0.3$ and $0.1\\pm0.2$ respectively. The combination of effective radius and surface brightness, $\\log r_e - \\beta \\log \\langle I \\rangle_e$, that enters the Fundamental Plane produces offsets smaller than $-0.02\\pm0.10$. This systematic error is insignificant and independent of galaxy magnitude or size. A catalog of photometry and surface brightness profile parameters is presented for three of the clusters in our sample, RX J0142.0+2131, RX J0152.7-1357, and RX J1226.9+3332 at redshifts 0.28, 0.83, and 0.89 respectively. ",
        "introduction": "Theoretical hierarchical models of galaxy formation predict that numerous small halos, primordial dwarf galaxies, collapsed early on in the history of the universe. These merged to form ever larger halos with deeper potentials and greater baryonic mass accumulation.  While these $\\Lambda$CDM models are successful at explaining structure formation on large scales, the agreement with observations is worse on the level of galaxies \\citep{con07}. Understanding the formation of primordial galaxies and how galaxies came to have the stellar populations, masses, and structural properties observed in the local universe is fundamental to our understanding of cosmology.  Generally, the physical processes a galaxy undergoes during its evolution depend primarily on only two factors: mass and environment.  While mass may affect galaxy structural and dynamical properties and star formation activity in predictable ways, the local environment of the galaxy is the source of complex evolutionary processes such as galaxy-galaxy merging, harassment, and stripping. Despite this, tight scaling relations of galaxy properties exist and must be reconciled with potentially different evolutionary paths.  Such scaling relations include the Tully-Fisher relation between disk galaxy total magnitude and rotational velocity \\citep{tf77}, the three-parameter Fundamental Plane (FP) relationship between early type galaxy velocity dispersion, effective radius, and effective surface brightness \\citep[e.g.][]{dlbd87,dd87}, the red sequence for early type galaxies, and relationships between absorption line strengths and central velocity dispersion. How these tight relations evolve with redshift or vary with galaxy mass places strong constraints on galaxy formation and evolution models.   Observational studies of cluster galaxies have found that the most massive galaxies are composed mainly of old stellar populations suggesting that massive early type galaxies formed at redshifts $> 2$ \\citep[e.g.][]{ble92,jorg99,tfwg00,blake03,mei06}. There is little evidence for recent star formation in these quiescent galaxies. They form a red sequence with little scatter in a color-magnitude diagram \\citep[e.g.][]{SV78} and they obey tight scaling relations in their kinematic and structural properties \\citep[e.g.][]{dlbd87,jfk96}.   These all suggest an early epoch of star formation in a homogeneously old galaxy population.  Although massive galaxies exhibit pure passive evolution since z $\\sim2$ evolving only in luminosity and color as stellar populations age, observations have also shown that clusters and lower mass cluster galaxies continue to evolve at intermediate redshifts.  The fraction of blue, star forming galaxies increases with redshift \\citep[e.g.][]{BO84, Ellingson01}  while larger fractions of spiral to S0 galaxies are observed in clusters at intermediate redshifts \\citep[e.g.][]{Dressler97}.  Meanwhile hierarchical structure formation models predict that galaxies form from the collective mergers of smaller galaxies. These mergers are expected to continue with high frequency through intermediate redshifts \\citep[e.g.][]{Baugh96}.  Semi-analytical models predict a characteristic mass a factor of 3 lower than found at z = 1 \\citep{poggianti04} with simulations underpredicting the numbers and mass densities of the most massive galaxies by a factor of over 100 \\citep{con07}, all implying that massive galaxies undergo a faster build-up than predicted. Dynamical models also often show a greater range in early type galaxy structure and kinematics than observed \\citep{deZF91}.  In order to reconcile apparently contradictory observations with formation and evolution models, we must study galaxy populations over a wide range of redshifts and for a wide range in galaxy mass. Previous studies have tended to investigate large samples of galaxies within a single cluster or epoch \\citep[e.g.][]{fz05, FPz0}, or fewer galaxies over a range of redshifts \\citep[e.g.][]{tsc02}, and often only the brightest, most massive galaxies.  Through the Gemini/HST Galaxy Cluster Project, we seek to better understand the processes driving cluster galaxy evolution by studying scaling relations as a function of mass and environment since z $= 1.0$, at a time when the universe was only half its current age.  We have therefore obtained data for a large sample of galaxies, which includes a sufficient number of galaxies at each redshift to accurately measure scaling relations and a sufficient number of redshifts to measure trends in the scaling relations as a function of redshift \\citep{jorg06,jorg07}.  At each redshift our sample spans a wide range in galaxy luminosity in order to investigate the role of galaxy mass. To discriminate between evolutionary effects due to environment and due to galaxy mass we have obtained consistent radial coverage for each cluster.   In a previous paper we have described the project and sample \\citep{jorg05}. Our sample of galaxy clusters consists of 15 X-ray selected massive clusters with redshifts ranging from $0.15 < z < 1$. For each of the clusters in our sample, we have obtained Gemini/GMOS spectra for $\\sim 30$ cluster members with a wide range of luminosities. The targetted objects are chosen independently of morphology since there is evidence that morphologies may continue to evolve even to the present epoch and we wish to avoid ``progenitor bias\" \\citep{vdf01}. From the spectra, we obtain redshifts, velocity dispersions, and line index measurements.  Using HST imaging data, we measure structural parameters and surface brightness profiles. With this large sample, we are studying scaling relations such as line index strengths and the FP as a function of redshift.   Line index measurements probe chemical enrichment providing insight into the star formation histories of the galaxies. The FP is examined in terms of mass and mass-to-light (M/L) ratios in order to probe both assembly histories and luminosity evolution.   We are using structural parameters for quantitative measurements of morphology to study morphological evolution. We have published the results on line index relations for RX J0152.7-1357 \\citep{jorg05} and RX J0142.0+2131 \\citep{barr05}, and the FP for RX J0142.0+2131 \\citep{barr06} and RX J0152.7-1357 and RX J1226.9+3332 \\citep{jorg06,jorg07}.   The FP relates surface brightness, effective radius, and velocity dispersion in a tight and well established relation \\citep{dlbd87,dd87,jfk96} \\begin{equation} \\log r_{e} = \\alpha \\log\\sigma + \\beta \\log\\langle I\\rangle_{e} + \\gamma. \\end{equation} It can also be described as a relation between M/L and $\\sigma$ or as \\begin{equation} \\log (M/L) = \\xi \\log M + \\gamma^{\\prime} \\end{equation} which makes the FP a powerful tool to study the star formation and assembly history in early type galaxies. A comparison of the FP for our high redshift sample to that of low redshift Coma cluster galaxies has revealed a number of interesting results.  The two z $\\sim 0.85$ clusters exhibit a steeper slope than the low redshift FP. We find this to be evidence for downsizing in which the lower mass galaxies have undergone more recent star formation and are overluminous compared to their low-z counterparts \\citep{jorg06,jorg07}. The FP for RX J0142.0+2131 at z=0.28, while having the same slope as the low z sample, displays a greater scatter, possible evidence for the galaxies having undergone rapid bursts of star formation during a cluster merger at z $> 0.85$ \\citep{barr06}.  Future papers will address the FP, line index scaling relations, and galaxy quantitative morphologies for all clusters in our sample in order to compare the star formation and assembly histories over a range of redshifts. Here we present the data reduction and analysis techniques we have used to measure structural and photometric parameters for three clusters in our sample, RX J0152.7-1357, RX J1226.9+3332, and RX J0142.0+2131.  RX J0152.7-1357 is a massive cluster at z=0.83 originally discovered from ROSAT data.  XMM-Newton and Chandra observations later showed the cluster to consist of two subclumps in the early stages of merging \\citep{Jones04,Maughan03,Girardi05}.   Much research has been done on this cluster, including recent studies of star formation rates \\citep{homeier05}, morphology, and the color-magnitude diagram \\citep{blake}. RX J1226.9+3332 is a massive cluster at z=0.89 and also the most X-ray luminous cluster known at such high redshift.  It was discovered  in the Wide Angle ROSAT Pointed Survey \\citep{ebe01} and exhibits a relaxed morphology. However, in a deep XXM-Newton and Chandra study of the X-ray mass analysis of this cluster, \\cite{MJJV07} find evidence for a recent or ongoing merger event. RX J0142.0+2131, at z=0.28, was first identified as a massive cluster in both the Northern ROSAT All-Sky Galaxy Cluster Survey \\citep{BVH00} and the ROSAT extended Brightest Cluster Sample \\citep{Ebe00}.   Although it is a relatively poor cluster, it displays a large cluster velocity dispersion, yet shows no signs of substructure \\citep{barr05}.  In this paper we describe our measurements of the physical properties of the galaxies in our sample and analyze how uncertainties affect the derivation of the scaling relations and overall results of our study. In Sections 2 and 3 we describe our observations and our imaging data reduction process. Catalogs of photometric and structural parameters for RX J0152.7-1357, RX J1226.9+3332, and RX J0142.0+2131 are presented in Section 4.  We discuss internal and external consistency of our results and implications on the measurements of the FP in Section 5. A summary is presented in Section 6. We assume $H_{0} = 70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{m} = 0.3$, and $\\Omega_{\\Lambda} = 0.7$. ",
        "conclusions": "\\label{tests}  \\subsection{Internal consistency}\\label{models}  Simulated galaxies are used to investigate sources of error in our pipeline and test the accuracy of the software used. GALFIT, for example, produces uncertainties for the output parameters which tend to be too low since these are strictly random uncertainties and do not take into account the fact that the profile we fit each galaxy with may not be an exact match to the true galaxy profile \\citep{peng01}.  We therefore use the simulated galaxies to determine more realistic measurement uncertainties and to test how these uncertainties affect our FP measurements.  The simulated galaxies are also used to test the effect of different PSF models and sizes for convolution in the galaxy profile fitting.  In this section, we describe the methods used to simulate the galaxies, describe tests to determine the best PSF to use for the 2-D surface brightness profile fitting, and determine expected uncertainties in our galaxy parameter measurements.  \\subsubsection{Galaxy simulations}  To generate realistic galaxies, we used the structural parameters of 148 galaxies in our low redshift Coma cluster comparison sample \\citep{jf94}. We transformed the values of M$_B$ to the filter and redshift of each cluster using stellar population models of \\citet{bc03}. For simulating galaxies at z $= 0.83$, we transform the M$_B$ magnitudes of Coma cluster galaxies to the observed $i^{\\prime}$ band using \\begin{equation} M_B = B_{rest} - DM(z), \\end{equation} \\begin{equation} B_{rest} = i^{\\prime} + 0.8026 - 0.4268(i^{\\prime} - z^{\\prime}) - 0.0941(i^{\\prime} - z^{\\prime})^2, \\end{equation} \\citep{jorg05,jorg07} and a mean empirical color, $(i^{\\prime} - z^{\\prime}) = 0.8$, for our early type galaxies at this redshift. For z $= 0.28$, we use \\begin{equation} B_{rest} = i^{\\prime} + 0.4753 + 1.6421(r^{\\prime} - i^{\\prime}) - 0.0253(r^{\\prime} - i^{\\prime})^2 \\end{equation} \\citep{barr05} with a mean empirical color $(r^{\\prime} - i^{\\prime}) = 0.5$. We do not include a separate simulation for z $= 0.89$ as those made for z $= 0.83$ are considered representative. We scale r$_{eff}$ to the appropriate redshift within our assumed cosmology. To create a large sample of simulated galaxies we added a small amount of random scatter to the magnitudes, r$_{eff}$, and Coma galaxy axial ratios while randomly generating positions and position angles. The Sersic parameter, n, was allowed to vary randomly between 0.5 - 5.5, except for a few sets of simulations where values ranged between 3.5 - 4.5 in order to increase the total number of early type galaxy simulations. The diskiness/boxiness of each galaxy was also allowed to vary randomly. Galaxies were added as perfect Sersic profiles, albeit with Poisson noise added to the galaxy images.  In order to simulate galaxies with multiple components, we also generated a set of bulge + disk galaxies.  In these cases, bulges were created by modeling from the Coma galaxy structural properties.  Disks were then added to the same position assuming disk-to-total light ratios ranging randomly from 0.25 - 0.65 with disk sizes ranging from (1 - 1.4) r$_{e(bulge)}$. The inclination was varied from 0 to 90 degrees, uniformly in $\\cos i$. As in \\citet{jf94}, we assign each component an intrinsic axial ratio: b/a(bulge) = 0.7 and b/a(disk) = 0.15. These values are chosen such that when transformed into observed values according to the inclination by \\begin{equation} b/a_{observed} = \\sqrt{1 - (1-(b/a_{intrinsic})^2)(\\cos{i})^2} \\end{equation} \\citep{sfs70}, they will display the same range of axial ratios found in real galaxies.  Simulated galaxies were added to real images using two different means. Galaxies were generated with ARTDATA by defining a Sersic profile and using the Coma based parameters.  They were also created using GALFIT itself by turning off the fitting for all parameters thereby forcing GALFIT to output models of the input parameters. No differences were observed in the galaxies generated or in the results obtained from the two methods, but both were used to ensure that no bias was produced from the simulation method.  Before being added to images, simulated galaxies were first convolved with a PSF and noise was added to the galaxies with the IRAF task MKNOISE.  PSFs were created using 4-6 unsaturated real stars having high S/N using the routines in the IRAF package DAOPHOT. These were created for each ACS image that simulated galaxies were added to.  In order to measure realistic GALFIT parameter uncertainties we generated a total of over $2600$ z $= 0.83$ and $2500$ z $= 0.28$ simulated galaxies. So as to not increase the surface density of objects in the images and thereby affect recovery of the surface brightness profile parameters, no more than 50 were added at a time to the real cluster images. It is important to note that all galaxies are generated as perfect Sersic profiles and recovered as Sersic and r$^{1/4}$ law profiles.  We do not simulate structure, such as arms, bars, double nuclei, shells, or twisted isophotes, all of which would effect the fitting. For the measurement of the cluster FPs, we have included only non-emission line, non-star forming galaxies which we take to be the early types. However, at higher redshift, one might expect to find more late-type features in these non-star forming galaxies. As a simple test of how added structure will affect single profile model fits, we use the set of simulated galaxies having both bulge + disk components.  In Figure \\ref{fprange}, we display a face-on view of the FP with our Coma cluster and high redshift galaxies plotted. We overlay our high redshift simulated galaxy sample.  Expected velocity dispersions given the input r$_{e}$ and $\\langle \\mu\\rangle_e$ for each simulated galaxy were calculated using the FP equation for Coma, $\\sigma = (\\log r_{e} + 0.82 \\log \\langle I\\rangle_{e} + 0.443)/1.3$ \\citep{jorg06}. It can be seen that our simulated galaxies span the same region of the FP as the real galaxies.  \\subsubsection{Choosing the best PSF for profile fitting}  In order to determine the most appropriate PSF to use for deconvolution during the surface brightness profile fitting, we use a subset of 700 simulated galaxies to test the recovery of parameters with GALFIT using several different PSF models and sizes. Images were run through our reduction pipeline starting with SExtractor object detection in order to provide initial guesses for the GALFIT profile fitting. We tested the GALFIT recovery using four different PSFs: the PSF based on real stars, a raw $3^{\\prime\\prime}$ Tiny Tim PSF generated for the appropriate location on the ACS chip, and $3^{\\prime\\prime}$ and $9^{\\prime\\prime}$ drizzled PSFs (see Section \\ref{fitting} for details). Galaxies were fitted with both Sersic and r$^{1/4}$ law profiles.  The r$^{1/4}$ law fits show greater discrepancy from input values than Sersic function fits because galaxies were added with a Sersic parameter n which varied randomly between 1 and 4.5 for these tests.  In Table \\ref{psfresults} we list the average differences and standard deviation between input and recovered magnitude, effective radius (in arcsec), and the combination of parameters that enter into the FP, $\\log r_e - \\beta \\log \\langle I \\rangle_e$, hereafter referred to as the Fundamental Plane parameter (FPP), where $\\langle I \\rangle_e$ is the surface brightness within $r_e$ in units of $L_\\odot$ pc$^{-2}$ and $\\beta \\sim -0.8$ \\citep{jorg06}. The average difference between input and recovered values is nearly identical regardless of the PSF used. However, we do find that using the raw Tiny Tim PSF consistently provides the worst results. Using $9^{\\prime\\prime}$ drizzled PSFs showed very slight improvement over $3^{\\prime\\prime}$ drizzled PSFs and very similar results to the real PSFs. A test using a subsampled PSF showed no improvement in the output results. Due to the difficulties producing a real PSF from the few stars in our images, we therefore chose to use the model $9^{\\prime\\prime}$ drizzled PSF. In the event that the Tiny Tim modeled $9^{\\prime\\prime}$ drizzled PSF is not the best representation of the true PSF, we note that the differences in the FPP due to the PSF used are insignificantly small.   \\subsubsection{Structural parameter measurement uncertainties}  The full set of simulated galaxies were used to investigate realistic uncertainties in structural parameter measurements.  Average offsets (recovered - input values) and the associated rms scatter in these differences are provided in Tables \\ref{devresults} and \\ref{fppresults}. As with the PSF tests, simulated galaxies were recovered using the same methods as real galaxies. The $9^{\\prime\\prime}$ drizzled PSF described in the previous section is used for the surface brightness profile fitting of both real and simulated galaxies.  We compare the r$^{1/4}$ law recovered parameters from GALFIT with the input values for all simulated galaxies in Figure \\ref{alldevres}. We find small average systematic offsets and associated rms scatter from the true magnitude and $\\log$ r$_e$ (arcsec) of $-0.2\\pm0.4$ and $0.1\\pm0.2$ for the high z sample. For the lower redshift set, we find offsets of $-0.4\\pm0.4$ and $0.3\\pm0.3$ respectively. Average measurement offsets and rms scatter for the FPP are only $-0.01\\pm0.06$ and $-0.04\\pm0.10$ for the two samples.  This is for the full set of simulated galaxies including those that were created with exponential profiles (n $\\sim 1$) but which were recovered with n = 4, r$^{1/4}$ law profiles.  To obtain realistic uncertainties in our FP parameters, we must compare only the range of parameters from galaxies that went into constructing our FPs.  We plot a histogram of the Sersic n values measured by GALFIT for our real galaxy samples in Figure \\ref{nhist}. The shaded histogram displays only those galaxies used in our FP.  Galaxies which were not included in our FP either had emission line spectra, $\\log$ Mass $< 10.3$, or n $< 1.5$ \\citep{jorg06}.  We therefore repeat the comparison between input and recovered parameters in Figure \\ref{devres} for galaxies created with n $> 2.0$ profiles.  It can be seen from Table \\ref{devresults} that while magnitude and size errors have decreased, the systematic error and random uncertainty in the FPP of $-0.01\\pm0.05$ and $-0.02\\pm0.10$ for the two samples exhibit little change.  The larger rms scatter for the lower redshift sample is due primarily to a few outliers with failed fits.  In our real galaxy sample all failed fits are flagged and refitted or left out of the FP measurement. We therefore consider this measurement uncertainty to be an upper limit.  Measurement errors are similar when fitting with a Sersic function, allowing n to vary. Systematic offsets are slightly lower for the magnitude and r$_{e}$ measurements, but there is no significant improvement in the FPP over fits made with r$^{1/4}$ law profiles. In Figure \\ref{serres}, we display the error in recovered n as a function of input n.  It is apparent that random error increases with increasing n.  This is likely because fitting of n=1 type galaxies tends to be more robust than n=4 galaxies where there is greater weight in the wings of the galaxy and fits are more easily affected by the sky values and near neighbors.  Because we expect the intrinsic value of n to affect the results for r$^{1/4}$ law fits, we plot the error in r$^{1/4}$ law recovered parameters as a function of input Sersic n in Figure \\ref{deln}.  Fluxes and sizes are overestimated for intrinsically low n galaxies because the r$^{1/4}$ law profile imposes broader wings than the galaxy has.  Magnitude and half-light radius are conversely underestimated at high n ($> 4$). Because flux and size are highly correlated, the errors in the FPP are reduced, but can be as large as 0.05 for n = 1 exponential disk galaxies fitted incorrectly with a r$^{1/4}$ law profile. A linear fit to the offset from the true input values finds $\\Delta$FPP = 0.017 n - 0.067.  We investigate other factors such as near neighbors which might affect the fits. Simulated galaxies with close neighbors within 1 r$_e$ or within 1 arcsec display a greater deviation in the recovered vs. input values for the FPP than those galaxies with a nearest neighbor at a distance greater than 2 arcsec.  In Figure \\ref{delnbr}, we display the errors as a function of nearest neighbor distance. Points enclosed by triangles have neighbors within 1 r$_e$.  From the plot it is apparent that these points have a larger dispersion than those with more distant neighbors.  The dispersion is 3 times as large for objects having neighbors within 2 arcsec as compared to those with nearest neighbors $> 6$ arcsec. We also test whether errors change as a function of r$_e$, magnitude, or surface brightness (Table \\ref{fppresults}). We do find that random errors in the FPP may increase slightly at fainter magnitudes and smaller sizes. We find no trend as a function of the input c parameter, a measure of the galaxy core's intrinsic diskiness/boxiness.  Finally, we investigate how single profile fits are affected by multiple component, bulge + disk, galaxies.  Forcing a single r$^{1/4}$ law fit to a bulge embedded in a disk recovers a magnitude comparable to the sum of the two components and a half-light radius that is larger than for the bulge alone. When combined to form the FPP parameter, the average systematic offset from the FPP for the bulge component alone is $0.07$ with $\\sigma = 0.09$ (Table \\ref{devresults}). This offset is larger than for the case of single profile galaxies although similar to the random error in single profile cases. Obviously the extended exponential disk will affect the fit, but the resultant FPP is similar to that of the bulge FPP alone within the expected uncertainties. Thus, we expect a r$^{1/4}$ law fit to a distant galaxy with large bulge and low surface brightness disk will recover primarily the bulge component with only a slight offset in the FPP. Likewise, we would expect the spectroscopic measurement of the velocity dispersion to come largely from light within the bulge.  For cases where faint disks are not obvious, we expect FP measurements to primarily represent the bulge components. We conclude from our simulated galaxy analysis that while r$^{1/4}$ law fits may lead to systematic errors in individual parameters for real galaxies, we do not expect this to be a problem for our FP analysis. We explore this further in Section \\ref{ext}.  \\subsubsection{Comparing real galaxy measurements}  For both RX J0152.7-1367 and RX J1226.9+3332, a number of galaxies were observed multiple times in overlap regions from different telescope pointings.  As one final test of our internal consistency we compare the derived structural and photometric properties from these multiple measurements. In Figure \\ref{intcomp}, we compare the difference in derived FPP for both Sersic and r$^{1/4}$ law profile fits in pairs of images.  Different symbols denote different image pairs and the corresponding dashed lines show the average difference for each pair. Average differences for pairs range from (absolute value) 0.0007 to 0.015 with standard deviations of 0.002 to 0.015. These differences are generally smaller than GALFIT measurement uncertainties, and the scatter in these real measurement differences is smaller than the random uncertainties determined from simulated galaxies.  We do not find any worrisome systematic trends between images; all differences between images are within expected measurement uncertainties.   \\subsection{External consistency}\\label{ext}  \\citet{blake} have used the same ACS dataset as used in this paper to study the structural parameters and colors of 149 galaxies in RX J0152.7-1357. They also perform surface brightness profile fitting using GALFIT with a procedure very similar to ours.  Images are processed in a similar manner, and like us, they fit all neighbors with $i_{F775W} < 25 AB$ while masking objects fainter than this. Differences in methodology are few but three significant differences exist.  While we use a $9^{\\prime\\prime}$ Tiny Tim model PSF drizzled in the same manner as our data, they choose to use empirical PSFs from archival HST images.  They also impose an upper limit for the fitted Sersic parameter, n, constraining the value to n $\\leq 4$, whereas we applied no constraints.  Because of the strong coupling between Sersic function parameters, this will cause differences in the measured values of both r$_{e}$ and total magnitude for galaxies we find with n $> 4$.  Finally, while we hold the sky value constant during the fitting at the SExtractor measured local value, they simultaneously fit for the sky value with GALFIT, except in cases where this led to poor fits.  We find a total of 26 objects in common, 3 of which have large discrepancies in magnitude. These were galaxies that were not well fit with single Sersic functions and which had late type features evident in the residual images.  Excluding these,  we compare the measured Sersic surface brightness profile parameters of the remaining 23 galaxies to the values listed in the Blakeslee et al. catalog. We plot the differences in our measurements in a $\\Delta \\log r_e$ vs. $\\Delta \\langle \\mu\\rangle_e$ plane (Figure \\ref{offsetbk}).  From a linear fit to these data points we measure a magnitude zeropoint shift of $-0.025$ between the two works from the offset in the intercept from zero.  This is in the sense that we find objects to be on average 0.025 mag brighter. We find that large differences in measured surface brightness or size are correlated with large discrepancies in the recovered Sersic n parameter (see also Figure \\ref{deln}). Galaxies with measured n differing by less than 0.1 for the two groups are circled. Correcting for the zeropoint shift we compare the difference in magnitudes, r$_{e}$, n, and the FPP. Results are listed in Table \\ref{bkresults} and displayed in Figure \\ref{bkcomp}.  The increasing difference in measured n at n $>4$ is easily explained by the constraints Blakeslee et al. impose on the Sersic n. Blakeslee et al. chose not to use Tiny Tim models finding that analytical representations for the PSF overestimated fitted r$_{e}$ values.  However, with our $9^{\\prime\\prime}$ model PSF we find r$_{e}$ values 2\\% smaller on average than their measurements. Average differences in total measured magnitude and effective radius are insignificantly small compared to the rms scatter. However, there appears to be a trend as a function of size and magnitude.  We find brighter galaxies to be brighter than found by Blakeslee et al. and fainter galaxies to be fainter.  Likewise, we measure larger sizes for larger galaxies and smaller sizes for smaller galaxies in comparison to Blakeslee et al.  This can be attributed in part to different values recovered for the Sersic n parameter.  In five cases where we find similar values within 0.1 for measured n, magnitude and size discrepancies are negligible regardless of size or magnitude. For 10 objects measured by Blakeslee et al. to have $n = 4.0$ due to their constraint on the Sersic index, and for which we find best Sersic function fits with $n > 4.0$,  we compare the Blakeslee et al. results to our r$^{1/4}$ law ($n = 4.0$) fits.  We find the average offsets in magnitude and size are similar but the random scatter is nearly half that of the differences in Sersic measurements for these same galaxies. For large galaxies, in cases where we find $n > 4.0$, the covariance of Sersic function parameters generates fits with larger $r_e$ than found by Blakeslee et al. For small galaxies which are more strongly affected by the choice of PSF, we speculate that our broad $9^{\\prime\\prime}$ PSF may be the cause of the smaller size measurements as compared to Blakeslee et al. Due to the coupling of Sersic function parameters, a smaller recovered size results in a fainter recovered magnitude. In Figure \\ref{offsetbk} we also show the parameter measurement errors in the $\\Delta \\langle \\mu \\rangle_e - \\Delta \\log r_e$ plane for our z $= 0.83$ simulations fitted with an r$^{1/4}$ law profile.  We find the same coupling of errors and an offset in magnitude of 0.018, similar to that found for the Blakeslee et al. comparison.  Since we used PSFs generated from real stars to create the simulated galaxies and the $9^{\\prime\\prime}$ model PSF to fit the galaxies, this further suggests that the broad PSF may be cause of the slight magnitude differences.  Because of the correlation between Sersic profile parameters, the FPP is insensitive to these small variations in size measurement.  When taking into account the small magnitude zero point differences, we find the differences between the Blakeslee et al. and our FPP measurement to be completely negligible and smaller than the systematic error of 0.006 determined from galaxy simulations.  The rms scatter in the differences between our measurements is a factor of 5 smaller than the uncertainty expected from the simulations.  The systematic difference in the FPP measurement of 0.009 is larger when not including the zero point shift, but is still much smaller than our expected random uncertainty of 0.05 determined from simulations.  \\subsection{Implications for FP measurements}\\label{discFP}  Although we have fit each galaxy with both Sersic and r$^{1/4}$ law functions, we use measurements of r$_{e}$ and surface brightness from the r$^{1/4}$ law fits to derive the FP since the galaxies in our low redshift comparison sample were originally fitted only with r$^{1/4}$ law profiles. Not all of our sample consists of pure r$^{1/4}$ law galaxies, and the Sersic function fits find a wide range of values for the Sersic parameter. However, only galaxies with n $> 2$ are used in our FP analysis \\citep{jorg06,jorg07,barr06}. In this section we explore how galaxy parameter measurement uncertainties translate to FP measurement uncertainties and whether any systematic biases exist in our methodology that would affect the derivation of the FP.  A plot of magnitude and FPP differences between GALFIT Sersic and r$^{1/4}$ law fits of real galaxies vs Sersic index, n, shows very good agreement for n $\\sim4$ galaxies as expected (Figure \\ref{mre}). Many galaxies were best fit with n $< 4$ and, for these, the r$^{1/4}$ law measured magnitude was as much as 1 mag brighter than that derived from the Sersic fits. Although output parameters from Sersic and r$^{1/4}$ law fits often differed, a plot of the difference in magnitude and in surface brightness vs. difference in log r$_e$ for the two fits shows the parameters to be strongly correlated and explains why differences in the FPP, although also correlated with Sersic n, are smaller. Galaxies with best Sersic fit $n > 2.0$ are displayed as larger points, but note that even galaxies with low Sersic index follow the same linear relation. Indeed, the FPP is reasonably robust against the form of the profile used to fit the galaxies.  This effect has been noted by \\citet{kidf00}, \\citet{tgc01}, and \\citet{l97}. While the uncertainties in measured structural and photometric parameters may seem large in cases where $n < 4.0$ according to best Sersic function fits, the FPP has a much lower $\\sim0.01$ systematic error with $0.1$ random uncertainty. We find that even when we fit a pure exponential disk ($n=1$) with an r$^{1/4}$ law profile, the errors in the magnitude and r$_{e}$ are large but the average error in the FPP is only 0.05.  In our study of the FP we investigate whether the slope changes as a function of redshift.  A change in slope would imply a change in M/L ratio as a function of galaxy mass.  It is therefore critical to ensure our measurements are not biased as a function of magnitude or size.  From our galaxy simulations we find some evidence for a slight increase of up to a factor of 3 in the scatter of the FPP at smaller sizes ($< 1$ arcsec) and at fainter magnitudes ($i > 22.5$).  However, this effect is small. More significantly, we find a strong correlation between the uncertainty in the FPP and the difference in measured vs. intrinsic Sersic index n.  Exponential disk galaxies with $n < 2$ fitted with r$^{1/4}$ law profiles have systematic offsets as large as the random uncertainty. If there is a trend in real galaxy magnitudes with index n, this could affect the FP slope measurement.  In Figure \\ref{nmtrend}, we display the measured Sersic index n versus SExtractor total magnitude for the galaxies in our RX J0152.7-1367 and RX J1226.9+3332 FP samples.  For the most part, the points exhibit random scatter. There is a hint of a trend in that the faintest galaxies have a slightly lower average Sersic index n.  The correlation coefficients for these two samples are small, only -0.31 and -0.14 for RX J0152.7-1367 and RX J1226.9+3332 respectively. Although this slight correlation coupled with the trend of increasing FPP measurement offset with decreasing n may affect the FP slope, this effect is much smaller than the random uncertainties in the FPP.  It is also important to ensure our measurement uncertainties are insensitive to redshift.  A redshift bias could be caused by, for example, GALFIT parameter measurement uncertainties being systematically different for objects with larger sizes or brighter magnitudes at low redshift as compared with smaller, fainter objects at high redshift. We do not find any evidence for this in our recovery of simulated galaxy parameters in simulations of the high (z = 0.83) and low (z = 0.28) redshift clusters.  The average magnitude and effective radius measurement uncertainties are slightly larger for the low redshift sample but scatter in the recovered values is about the same for both clusters.   The FPP has an average offset of recovered values from input which is insignificant for both populations, and a random uncertainty which is larger for the lower redshift cluster due to a small number of outliers with particularly poor fits. In our real galaxy sample all failed fits are flagged and refitted or left out of the FP measurement.  Through this simulated galaxy analysis we have found a number of other factors which affect the galaxy structural parameter measurements. Near neighbors affect profile fits and based on simulations, we find that the highest surface brightness, most compact galaxies are affected the least. However, the only effect this has on lower surface brightness galaxies is to increase the random uncertainty by up to factor of 4 for only those few galaxies which have a neighbor within 1 arcsec or 1 r$_{e}$. No systematic variation as a function of surface brightness is expected.  Although we cannot simulate galaxies with a full range of substructures, nor do we know the true distribution of such features in our real data, we estimate uncertainties for galaxies with more complex structures with bulge + disk simulations.  We find a systematic offset of $0.07$, much larger than the 0.01 for simulated galaxies generated with a single profile, with a slightly larger rms scatter of 0.10. With a large sample of galaxies exhibiting complex morphologies, we would expect measurements of the FP to be affected by the larger uncertainties. However, we check the goodness of all fits from residual images.  Any features not modeled by the simple Sersic/r$^{1/4}$ law profile fits will show up in these images as imperfect subtraction from the galaxy model. For most of our galaxies, subtraction is good without any obvious sign of extra arms, twisted isophotes, low surface brightness disks, or any other structures.  Galaxies with such features tend to be best fit with Sersic $n < 2.0$, and these objects are not included in our FP measurements.  In conclusion, we take final uncertainties in the FPP to be $\\sim 0.1$ based on our simulations, dominated by random uncertainties. We find no systematic offsets between recovered and input values for the FPP as a function of magnitude or half-light radius from our galaxy simulations. While we do find uncertainties are affected by near neighbors or Sersic index n, we do not expect the shape or tilt of the FP to be affected by these uncertainties, unless for example all lower mass galaxies have exponential profiles, which we see no evidence for.   We find that our measurement uncertainties should not vary with redshift unless, for example, higher redshift galaxies contain more late-type morphological features.  Although this is a possibility, we would expect to spot this in our residual fitting images. We therefore do not find any evidence for measurement uncertainties or systematic biases introduced by our data reduction methods to affect our derived FP as a function of redshift, galaxy size, or magnitude.  We expect the measured tilts to be intrinsic values for each cluster."
    },
    "0910/0910.0661_arXiv.txt": {
        "abstract": "The large survey programs being performed nowadays, being the SDSS their flagship, provide us with morphological parameters which allow for extraction of large galaxy samples. We will analyze the methodology for obtaining an AMIGA-like catalogue of isolated galaxies from the SDSS~DR5 photometric catalogue of galaxy objects, together with the roadblocks found in the process, and suggested workarounds. ",
        "introduction": "% \\label{sec:introduction} The Catalogue of Isolated Galaxies (CIG) by \\citet{1973AISAO...8....3K}, basis for the AMIGA sample \\citep{2005A&A...436..443V}, was compiled using a visual search on photographic plates. In the era of the Virtual Observatory and large surveys, however, the automated compilation of isolated galaxy samples not only becomes feasible, but potentially less biased and error prone.  There are previous examples of automatic extraction of isolated galaxy samples in the literature, some of them mimicking Karachentseva's criterion, as is the case of \\citet{2005AJ....129.2062A}. The search of compact groups uses similar constraints, and a recent example of automated extraction of them in galaxy catalogues is shown by \\citet{2009MNRAS.395..255M}.  With automatically extracted samples, combined with the extra parameters derived by the SDSS pipeline, it is possible to estimate the degree of isolation of the selected galaxies, using for instance the tidal forces or local density estimations characterised in \\citet{2007A&A...470..505V,2007A&A...472..121V}. Additionally, we gain the ability to re-do the studies, and apply them to similar catalogues in the southern hemisphere, so that an all-sky sample of isolated galaxies can be finally assembled. In this proceeding we will focus on the automatic extraction of a CIG (or AMIGA) like sample of isolated galaxies. ",
        "conclusions": "% \\label{sec:conclusions}  We have seen that in order to use the SDSS Photometric catalogue for retrieving galaxy properties, and estimate their isolation degree, the basic star-galaxy separation provided by the SDSS pipeline is not enough, and more work is needed in order to be able to obtain meaningful results regarding the isolation of galaxies.  The analysis of available parameters, however, has shown to be promising in achieving a more accurate star-galaxy separation which, combined with data mining and clustering tools, may finally allow a reuse of SDSS photometric data for this kind of studies.  In the future, we will continue pursuing an automated isolated galaxy sample builder system, using data mining techniques in order to find out the parameters with the best star-galaxy separation properties, and be able to extend the selection of isolated galaxies to similar catalogues in the future."
    },
    "0910/0910.2178_arXiv.txt": {
        "abstract": "\\small The methods of determining the fractal dimension and irregularity scale in simulated galaxy catalogs and the application of these methods to the data of the 2dF and 6dF catalogs are analyzed. Correlation methods are shown to be correctly applicable to fractal structures only at the scale lengths from several average distances between the galaxies, and up to (10--20)\\% of the radius of the largest sphere that fits completely inside the sample domain. Earlier the correlation methods were believed to be applicable up to the entire radius of the sphere and the researchers did not take the above restriction into account while finding the scale length corresponding to the transition to a uniform distribution. When an empirical formula is applied for approximating the radial distributions in the samples confined by the limiting apparent magnitude, the deviation of the true radial distribution from the approximating formula (but not the parameters of the best approximation) correlate with fractal dimension. An analysis of the 2dF catalog yields a fractal dimension of $2.20\\pm0.25$ on scale lengths from 2 to 20 Mpc, whereas no conclusive estimates can be derived by applying the conditional density method for larger scales due to the inherent biases of the method. An analysis of the radial distributions of galaxies in the 2dF and 6dF catalogs revealed significant irregularities on scale lengths of up to 70 Mpc. The magnitudes and sizes of these irregularities are consistent with the fractal dimension estimate of $D ={}$2.1--2.4\\par ",
        "introduction": "The spatial distribution of galaxies bears signatures of both the initial conditions in the early Universe and the evolution of the primordial density perturbations. An analysis of various galaxy samples performed using the two-point correlation function showed that this function has a power-law form $\\xi(r)=(r_0/r)^{\\gamma}$ on scale lengths ranging from 0.01 to 10 Mpc (hereafter we adopt a Hubble constant of $H_0 = 100 \\mathrm{ km/s/Mpc}$) with a slope of $\\gamma=1.77$ and the parameter $r_0=5$ $\\mathrm{Mpc}$ \\cite{Peebles}. It has long been considered that the scale of the $r_0$ parameter is the typical irregularity scale length, and the distribution of galaxies becomes uniform starting from the scale length of $r_0 = 5 $ Mpc. However, the discovery of structures with the scale lengths of several tens and hundreds Mpc \\cite{Baryshev} in recent surveys has cast doubt upon this hypothesis.  In this context, the problems of applicability limits and reliability of the correlation methods of the analysis of spatial distribution of galaxies, and finding new methods for describing large and very large structures acquire special importance.  At present, two kinds of data on the galaxy redshifts are of great importance.  \\begin{itemize} \\item The first kind are the redshift catalogs covering large areas (solid angles) of the sky, but limited to small redshifts (up to $z\\lesssim 0.5$) (2dF, 6dF, SDSS, etc.). Such catalogs can be analyzed via applying the correlation methods to determine the fractal dimension.  \\item The second kind is represented by the deepfield catalogs of photometric redshifts. Such studies cover small solid angles (of the order of $1^\\circ \\times 1^\\circ$), but extend to much larger redshifts $z>1$ (up to~6) (COSMOS, HDF, HUDF, FDF and others). Correlation methods are difficult to apply to such catalogs due to the small radius of the largest sphere that fits entirely inside the small solid angle considered. \\end{itemize}  However, both kinds of catalogs can be used to analyze the radial distribution of galaxies, built upon a sample confined by the limiting apparent magnitude. This method not only removes the restriction on the size of the largest sphere thereby significantly increasing the attainable research scale lengths, but it can also be applied to all galaxies in the catalog and not only to those in a volume-limited sample thereby increasing the number of objects studied. An analysis of fluctuations in the radial distribution of galaxies can be used to determine both the sizes and the amplitudes of the largest structures in the galaxy sample considered.  In this paper we analyze two methods of statistical analysis of structures---a determination of the fractal dimension, and an analysis of radial distributions. Despite the fact that our analysis is limited to the 2dF and 6dF catalogs, we constructed our simulated lists with two kinds of catalogs (covering large and small solid angles on the sky).  In this paper we make use of our own software, developed to simulate three-dimensional catalogs of galaxies and to perform statistical analysis of both real and simulated samples. It is a C++ library of functions (so far, without a user interface). We are currently preparing its description, which will be made available, along with the source code, at our web site. The software covers a somewhat broader scope of problems than that described in this paper, and will be a basis for a future package meant for comprehensive statistical analysis of the spatial distribution of galaxies. ",
        "conclusions": "As a conclusion, we shall specify the following results of our analysis of simulated catalogs of the spatial distribution of galaxies, and absolute magnitude distribution of galaxies: \\begin{itemize} \\item correlation methods can be correctly applied only on scale lengths from several average distances between the galaxies up to (10--20)\\% of the radius of the largest sphere that fits entirely inside the set. The authors of earlier studies believed that the method could be applied out to the entire radius, and the above 10\\% restriction, which applies to all distributions but uniform, was not taken into account when determining the scale length where the distribution becomes uniform (see, e.g., \\cite{SL}, \\cite{Tikhonov}) \\item the empirical formula (\\ref{efrr}), which is often used to approximate radial distributions of objects in magnitude-limited catalogs, yields equally adequate minimum root-mean-square approximation for both uniform and fractal distributions with dimensions exceeding 2.0. At smaller dimensions the scatter becomes too large and the formula is inapplicable.  We found the fractal dimension to correlate with the deviation of the true radial distribution from the approximating formula, and not with the parameters of the best-fit approximation. \\end{itemize}  Our analysis of real catalogs yielded the following results:  \\begin{itemize} \\item the data of the 2dF catalog imply a fractal dimension of $2.20 \\pm 0.25$ in the interval from 2 to 20 Mpc. No reliable conclusions can be made on larger scales about the dimension and scale of irregularities due to the intrinsic biases of the method. Deeper surveys and surveys with better sky coverage are needed for this task.  Because of its incompleteness, the 6dF catalog can not yet be used to derive a reliable estimate for the fractal dimension;  \\item An analysis of radial distributions revealed the significant irregularities both in the 2dF and 6dF catalogs. Deviations from smooth distribution exceed $7\\sigma_p$ and their scale lengths amount to 70 Mpc. The scale length and magnitude of irregularities correlate rather well with fractal-dimension estimates in the 2.1--2.4 interval. \\end{itemize}"
    },
    "0910/0910.2452_arXiv.txt": {
        "abstract": "We present results of a multi-wavelength study of episodic plasma injection into the corona of AR 10942. We exploit long-exposure images of the \\hinode{} and Transition Region and Coronal Explorer (TRACE) spacecraft to study the properties of faint, episodic, ``blobs'' of plasma that are propelled upward along coronal loops that are rooted in the AR plage. We find that the source location and characteristic velocities of these episodic upflow events match those expected from recent spectroscopic observations of faint coronal upflows that are associated with upper chromospheric activity, in the form of highly dynamic spicules. The analysis presented ties together observations from coronal and chromospheric spectrographs and imagers, providing more evidence of the connection of discrete coronal mass heating and injection events with their source, dynamic spicules, in the chromosphere. ",
        "introduction": "Observing the mechanics of the process that raises the 10,000 K chromospheric plasma to several million K has eluded the community since the inference of hot coronal plasma was made at the dawn of the rocket age \\citep{Edlen1943}. Recently, analysis of data from the \\hinode{} spacecraft \\citep[][]{Hinode} has established the connection of chromospheric and coronal heating processes \\citep{DePontieu2009}. The key to this observational-driven leap in understanding is the discovery of a second class of ``spicule'' \\citep[e.g.,][]{Roberts1945,Beckers1968}. These so-called type-II spicules were revealed in the high spatial resolution observations of the chromospheric limb provided by the Solar Optical Telescope \\citep[][]{Tsuneta2008}. \\citet{DePontieu2007b} showed that they originate in strong magnetic regions, are longer (4-8Mm), and display larger upward velocities (50-150km/s) than their classical (shock-driven) ``Type-I'' counterparts \\citep[2-5Mm tall with speeds 10-40km/s;][]{DePontieu2004,Hansteen2006,DePontieu2007a,Rouppe2007}.  \\citet{DePontieu2009} demonstrated the spatio-temporal correlation of upper chromospheric activity, in the form of Type-II spicules, and a weak component (2-5\\% of the peak line emission) in the blue wings of three emission lines formed in the transition region and corona (\\ion{C}{4} 1548\\AA{} at 150,000K, \\ion{Ne}{8} 770\\AA{} at 600,000K, and \\ion{Fe}{14} 274\\AA{} at 2MK). This correlation was strongest in locations of unipolar plage, as well as the supergranular network vertices of the quiet sun and coronal holes. In each case the blue-wing component, indicative of the presence of an upward moving jet, shows a nearly identical velocity distribution to that of the dynamic spicules seen in the same unipolar regions. This led \\citet{DePontieu2009} \\citep[and][]{McIntosh2009a} to suggest that these episodic heating events rooted in chromospheric activity play a significant role in filling the Sun's upper atmosphere with hot plasma.  \\begin{figure} \\epsscale{1.15} \\plotone{f1} \\plotone{f1b} \\caption{Context imaging and spectroscopy of AR 10942 on 2007, February 20 from XRT (top left), \\trace{} 1600\\AA{} (top right) and EIS in the \\ion{Fe}{12} 195\\AA{} (bottom left) and \\ion{Fe}{14} 264\\AA{} (bottom right) lines. We show the EIS and different SOT fields of view for a short-exposure co-temporal sequence (dot\\--dashed line) and for a longer exposure sequence taken beforehand (dot\\--dot\\--dashed line), and the reference position used to compute the space-time (x-t) plots used throughout this Letter. The bright plage and network emission in the TRACE image are outlined by an intensity contour of 250~DN. \\label{f1}} \\end{figure}  In this Letter we explore a long-exposure observing sequence performed by \\hinode{} and the Transition Region and Coronal Explorer \\citep[\\trace;][]{Handy1999} of solar active region (AR) 10942. Observations with the \\hinode{} X-Ray Telescope \\citep[XRT;][]{Golub2007} revealed the presence of episodic, fast, ``blobs'' of coronal material ($>1$MK) that appear to emanate from the plage in the active region \\citep[][]{Sakao2008}. Our motivation comes from the fact that the reported blob velocities (qualitatively) matched those of Type-II spicules and associated weak coronal upflows (measured spectroscopically) observed in the magnetic footpoints of AR coronal loops. In the sections that follow we explore the thermal connectivity and origins of these coronal blobs using an extended observational dataset, taken the day before the observations discussed by \\citet[][]{Sakao2008}. ",
        "conclusions": "We have shown that high velocity upflow events are visible in the chromospheric and coronal footpoints of loops that carry plasma blobs observed by \\trace{} and XRT. These velocities correspond spatially and spectrally to weak, blue-wing asymmetries that are observed across a range of temperatures in many EIS lines. We suggest that the weak upflows are related to episodes of dynamic spicule activity. Further, we suggest that the observed coronal outflows seen in long-exposure \\trace{} and XRT image sequences of the corona are the result of the heating of chromospheric material to coronal temperatures in dynamic Type-II spicules and, as such, may be the signature of discrete episodic coronal heating and mass injection events. The nature of these Type-II spicules, in tracing out the roots of the coronal magnetic topology and mass loading, suggests a more appropriate (or descriptive) name for these spicules: ``radices'' (singular: ``radix'' \\-- Latin for ``root'')."
    },
    "0910/0910.2980_arXiv.txt": {
        "abstract": "{} {Results of the long-term (11 years, from 1996 to 2006) H$\\alpha$ and H$\\beta$ line variations of the active galactic nucleus of NGC 4151 are presented.} {High quality spectra (S/N$>50$ and R$\\approx$8 \\ \\AA) of H$\\alpha$ and H$\\beta$ were investigated. We analyzed line profile variations during monitoring period. Comparing the line profiles of H$\\alpha$ and H$\\beta$, we studied different details (bumps, absorbtion features) in the line profiles. The variations of the different H$\\alpha$ and H$\\beta$ line profile segments have been investigated. Also, we analyzed the Balmer decrement for whole line and for line segments. } {We found that the line profiles were strongly changing during  the monitoring period, showing blue and red asymmetries. This indicates a complex BLR geometry of NGC 4151 with, at least, three kinematically distinct regions: one that contributes to the blue line wing, one to the line core and one to the red line wing. Such variation can be caused by an accelerating outflow starting very close to the black hole, where the red part may come from the region { closer to the black hole than the blue part, which is coming} from the region having the highest outflow velocities.} {Taking into account the fact that the BLR of NGC 4151 has a complex geometry (probably affected by an outflow) and that a portion of the broad line emission seems to have not a pure photoionization origin, one can ask the question whether the study of the BLR by reverberation mapping may be valid in the case of this galaxy.} ",
        "introduction": "In spite of many papers  devoted to the physical properties (physics and geometry, see e.g. Sulentic et al. 2000) of the broad line region (BLR) in active galactic nuclei (AGN), the true nature of the BLR is not well known. The broad emission lines, their shapes and intensities can give us many information about the BLR geometry and physics. In the first instance, the change in the line profiles and intensities could be used for investigating the BLR nature. It is often assumed that variations of line profiles on long time scales are caused by dynamic evolution of the BLR gas, and on short time scales by reverberation effects (Sergeev et al. 2001). In a number of papers it was shown that individual segments in the line profiles change independently on both long and short time scales (e.g., Wanders and Peterson 1996; Kollatschny and Dietrich 1997; Newman et al. 1997; Sergeev et al. 1999). Moreover, the broad line shapes may reveal some information about the kinematics and structure of the BLR (see Popovi\\'c et al. 2004).   One of the most famous and best studied Seyfert galaxies is NGC 4151 (see e.g. Ulrich 2000; Sergeev et al. 2001; Lyuty 2005, Shapovalova et al. 2008 - Paper I, and reference therein). This galaxy, and its nucleus, has been studied extensively at all wavelengths.   The reverberation investigation indicates a small BLR size in the center of NGC 4151 (see e.g. Peterson and Cota 1988: $6\\pm4$ l.d.; Clavel et al. 1990: $4\\pm3$ l.d.; Maoz et al. 1991: $9\\pm2$ l.d.; Bentz et al. 2006: $6.6_{+1.1}^{-0.8}$ l.d.).  Spectra of NGC 4151 show a P Cygni Balmer and He I absorption with an outflow velocity from -500 ${\\rm km \\ s^{-1}}$ to -2000 $\\ {\\rm km \\ s^{-1}}$, changing with the nuclear flux (Anderson and Kraft 1969; Anderson 1974; Sergeev et al. 2001; Hutchings et al. 2002). This material is moving outward along the line of sight, and may be located anywhere outside $\\sim$15 light-days (Ulrich and Horne 1996).  An outflow is also seen in higher velocity emission-line clouds near the nucleus (Hutchings et al. 1999), such as multiple shifted absorption  lines in C IV and other UV resonance lines (Weymann et al. 1997; Crenshaw et al. 2000), while warm absorbers are detected in X-ray data (e.g. Schurch and Warwick 2002).  Some authors assumed that a variable absorption is responsible, at least partially, for the observed continuum variability of AGN (Collin-Souffrin et al.1996; Boller et al. 1997; Brandt et al. 1999; Abrassart and Czerny 2000; Risaliti, Elvius, Nicastro 2002). Czerny et al. (2003) considered that most of variations are intrinsic to the source, though a variable absorption cannot be quite excluded.  The nucleus of NGC 4151 emits also in the radio range. The radio image reveals a 0.2 pc two-sided base to the well-known arc-second radio jet (Ulvestad et al. 2005). The apparent speeds of the jet components relative to the radio AGN are less than 0.050c and less than 0.028c at  nuclear distances of 0.16 and 6.8 pc, respectively. These are the lowest speed limit yet found in a Seyfert galaxy and indicates non-relativistic jet motions, possibly due to thermal plasma, on a scale only an order of magnitude larger than the BLR (Ulvestad et al. 2005).  The observed evolution of the line profiles of the Balmer lines of NGC 4151 was studied by Sergeev et al.(2001) in 1988--1998 and was well modeled  within the framework  of the two-component model, where two variable components with fixed line profiles (double-peaked and single-peaked) were used.   Although the AGN of NGC 4151 has been much observed and discussed, there are still several questions concerning the BLR kinematics (disk, jets or more complex BLR) and dimensions of the innermost region. On the other hand, as was mentioned above, multi-wavelength observations suggest both the presence of an accretion disk (with a high inclination) and an outflow emission (absorption). Consequently, further investigations of the NGC 4151 nucleus are needed in order to constraint the kinematics, dimensions and geometry of its BLR.  This work is subsequent to Paper I with the aim of studying the variations of both the integrated profiles of the broad emission lines and segments along the line profiles, during the (11-year) period of monitoring of NGC 4151.  The paper is organized as follows: in \\S2 observations and data reduction are presented. In \\S3 we study the averaged spectral line profiles (over years, months and Periods I--III, see Paper I) of H$\\alpha$ and H$\\beta$, the line asymmetries and their FWHM variations, light curves of different line segments and the line--segment to line--segment flux and continuum--line--segment relations. In \\S4 we analyze the Balmer decrements. The results are discussed in \\S6 and in \\S7 we outline our conclusions. ",
        "conclusions": "This work is subsequent to the Paper I (Shapovalova et al. 2008) and is dedicated to a detailed analysis of the broad H$\\alpha$ and H$\\beta$ line profile variations during the 11-year period of monitoring (1996--2006). From this study (Section 3) it follows that the BLR in NGC 4151 is complex, and that broad emission lines are the sums of components formed in different subsystems:  1) the first component is photoionized by the AGN continuum (far red line wings, $V_r\\sim +4000$ and $+5000 \\ \\rm km \\ s^{-1}$ and central (at $V_r=\\pm3500 \\ \\rm km \\ s^{-1}$) segments for continuum flux $F_{\\rm c}<7 \\times 10^{-14} \\rm erg cm^{-2} s^{-1} A^{-1}$). This region is the closest to the SMBH.  2) the second component is independent of changes of the AGN continuum (far blue line wings, $V_r\\sim-5000 \\ \\rm km \\ s^{-1}$; $-4000 \\ \\rm km \\ s^{-1}$). It is possibly generated by shocks initiated by an outflow.  3) the third component, where the central parts of lines ($V_r<4000 \\ \\rm km \\ s^{-1}$) are formed, in high fluxes $F_{\\rm c} > 7 \\times 10^{-14} \\rm erg \\ cm^{-2} s^{-1} A^{-1}$ is also independent of the AGN continuum (possibly, outflow and jet).    Finally, we can conclude that the BLR of NGC 4151 may not be purely photoionized, i.e. besides photoionization, there could be some internal shocks contributing to the broad lines  see also discussion in Paper I). At least there are three different subregions with different velocity field and probably different physical conditions, that produce complex variability in the broad lines of NGC 4151 and changes in the line profile (very often temporarily bumps).  Our investigations indicate that the reverberation might not be valid as a tool to determine the BLR size in this AGN and that this AGN is not perfect for this method. Consequently, the results for the BLR size of NGC 4151 (given in the introduction) should be taken with caution."
    },
    "0910/0910.3454_arXiv.txt": {
        "abstract": "Using the Infrared Spectrograph aboard the \\emph{Spitzer Space Telescope}, we observed multiple epochs of 11 actively accreting T~Tauri stars in the nearby Taurus--Auriga star forming region. In total, 88 low-resolution mid-infrared spectra were collected over 1.5 years in Cycles~2 and 3.  The results of this multi-epoch survey show that the 10~$\\mu$m silicate complex in the spectra of two sources --- DG~Tau and XZ~Tau --- undergoes significant variations with the silicate feature growing both weaker and stronger over month- and year-long timescales.  Shorter timescale variations on day- to week-long timescales were not detected within the measured flux errors.  The time resolution coverage of this data set is inadequate for determining if the variations are periodic.  Pure emission compositional models of the silicate complex in each epoch of the DG~Tau and XZ~Tau spectra provide poor fits to the observed silicate features.  These results agree with those of previous groups that attempted to fit only single-epoch observations of these sources.  Simple two-temperature, two-slab models with similar compositions successfully reproduce the observed variations in the silicate features.  These models hint at a self-absorption origin of the diminution of the silicate complex instead of a compositional change in the population of emitting dust grains.  We discuss several scenarios for producing such variability including disk shadowing, vertical mixing, variations in disk heating, and disk wind events associated with accretion outbursts. ",
        "introduction": "The spatial distribution, size, and composition of silicate dust grains observed in circumstellar disks surrounding young pre-main-sequence stars result from a variety of physical and chemical evolutionary processes.  As they play a key role in the evolution of circumstellar disks toward planetary systems, an array of observational techniques has been employed to study the characteristics of these dust grains (sub-millimeter, millimeter, mid-infrared, polarization, etc., \\cite[e.g.][]{wein1989,stro1989,beck1990,skru1990}.  Building upon much ground-based infrared photometry and imaging sensitive to the blackbody emission and scattered light from dust grains in circumstellar environments, the \\emph{Spitzer} Infrared Spectrograph (IRS) has amassed an extensive catalog of near- to mid-infrared spectral energy distributions (SEDs) of young sources in the nearest star forming regions \\citep{kess2005,furl2006,sarg2009,wats2009}.  Modeling of the SEDs of this large sample of T Tauri stars (TTS) ---  known to possess thick and actively accreting disks --- succeeds in defining the spatial and size distributions of the dust grains in these systems, potentially giving astronomers direct insight into the evolutionary status of these planetary building blocks.  For instance, SED models suggest the presence of `puffed-up' inner disk rims near the dust sublimation radius for some sources \\citep{natt2001,dull2001,muze2003,muze2004} and constrain the sizes of the optically thin inner disk holes of transitional objects finding some to be as large as 40~AU \\citep{stro1989,skru1990,calv2002,espa2007a,espa2007b}.  In addition to the information encoded in the SEDs, grain sizes and composition have been inferred from mid-infrared emission features formed by the super-heated dust grains located at small radii and in the upper atmospheres of these protoplanetary disks.  High-sensitivity, low-resolution spectral surveys of the 10~$\\mu$m silicate complex have provided the opportunity for in-depth analyses of compositional variations of the dust grains with the most detailed models including laboratory measured opacities for a variety of sizes and species \\citep[see e.g.][]{bouw2003,meeus2003,hond2006,sche2006,sarg2009}.  The parameters returned by these models include dust temperatures and mass fractions for different grain types, sizes, and structures (i.e., crystalline and amorphous, small and large, forsterite, enstatite, olivine).  Less detailed compositional analyses have been conducted by binning the 10~$\\mu$m silicate complex into various wavelength regions corresponding to the peak emission of specific dust grain species.  The integrated fluxes over these wavelength bins are used to compute ratios commonly referred to as crystalline indices that give an approximation of the relative abundances \\citep[see e.g.,][]{bouw2001,przy2003,vanb2003,kess2005,kess2006,wats2009}.  Several groups have published single-epoch \\emph{Spitzer} IRS surveys of TTS in which they note source-to-source morphological differences of the silicate feature \\citep{forr2004,kess2005,furl2006,sarg2009,wats2009}.  \\citet{forr2004} were the first to note a relationship between the crystallinity versus amorphous nature of the silicate emission feature with the evolutionary status of the dust grain population and the disk.  The presence of crystalline grains implies that the once amorphous interstellar grains have been annealed and cooled, consistent with heating in the inner regions of an evolving protoplanetary disk.  \\citet{kess2005} and \\citet{sarg2009} concluded the same, suggesting that a disk evolutionary sequence can explain the variety of shapes observed for the silicate features in their Short-Low (5.2-14.5~$\\mu$m) and Long-Low (14.0-38.0~$\\mu$m) IRS spectra of TTS.  The conclusions about the evolutionary status of the disks are based solely on single-epoch observations and do not account for the dynamic nature of the TTS sources, which are characterized by multi-wavelength variability on timescales as short as days.  Variations of the morphology of the silicate emission feature on the timescales of days, months, or even years would likely complicate such compositional studies of dust grain populations such as those leading to the development of a disk evolutionary sequence.  In order to test the validity of the conclusions drawn from the single-epoch studies, we used \\emph{Spitzer} IRS to conduct a spectroscopic variability study of the 10~$\\mu$m silicate emission feature in 11 actively accreting classical TTS found in the Taurus--Auriga star forming region.  In 2 of the 11 sources, DG~Tau and XZ~Tau, significant variability of the 10~$\\mu$m feature is observed.  The unexpected short timescale variations and potential scenarios for this behavior are presented in this Letter.  A long paper presenting more detailed modeling of these features and the spectra of the sources with no detectable variability is forthcoming (Leisenring et al., in preparation).  \\begin{deluxetable*}{crcccrrr} \\tabletypesize{\\scriptsize} \\tablecaption{Observations, Continuum and 10~$\\mu$m Feature Measurements \\label{tbl:1}} \\tablewidth{0pt} \\tablehead{ \\colhead{Star} &   \\colhead{AOR} &  \t\\multicolumn{2}{c}{Date Observed} &  \t\\colhead{SL1 \\& SL2} &   \\colhead{Mean Slit} &  \t\\colhead{$[6.0]-[13.5]$} &  \t\\colhead{Silicate Flux}  \\\\ \\colhead{Name} &   \\colhead{ID} &   \\colhead{Gregorian} &\t\\colhead{Julian} &\t\\colhead{Cycles $\\times$ Ramps} &   \\colhead{Offset ($^{\\prime\\prime}$)} &\t\\colhead{Color\\tablenotemark{a}} &\t\\colhead{(10$^{-14}$ W m$^{-2}$)\\tablenotemark{b}}} \\startdata DG Tau & 3530496 & 2004-03-02 & 2453067 & 2 $\\times$ 6 sec & -1.15 & 2.63 $\\pm$ 0.007 & 1.93 $\\pm$ 0.370 \\\\ \\nodata & 15110144 & 2005-10-09 & 2453653 & 8 $\\times$ 6 sec & -1.97 & 2.80 $\\pm$ 0.009 & 4.05 $\\pm$ 0.908 \\\\ \\nodata & 15115008 & 2005-10-15 & 2453659 & \\nodata & -2.15 & 2.87 $\\pm$ 0.009 & 3.99 $\\pm$ 0.866 \\\\ \\nodata & 15115264 & 2006-03-17 & 2453812 & \\nodata & 0.00 & 2.90 $\\pm$ 0.011 & 6.24 $\\pm$ 0.076 \\\\ \\nodata & 15115520 & 2006-03-22 & 2453817 & \\nodata & 0.00 & 2.90 $\\pm$ 0.011 & 6.00 $\\pm$ 0.077 \\\\ \\nodata & 19488000 & 2006-10-17 & 2454026 & \\nodata & -0.23 & 2.66 $\\pm$ 0.010 & 5.30 $\\pm$ 0.057 \\\\ \\nodata & 19487744 & 2006-10-22 & 2454031 & \\nodata & -0.40 & 2.64 $\\pm$ 0.008 & 3.98 $\\pm$ 0.166 \\\\ \\nodata & 19487488 & 2007-03-21 & 2454181 & \\nodata & 0.23 & 2.75 $\\pm$ 0.010 & 5.78 $\\pm$ 0.102 \\\\ \\nodata & 19487232 & 2007-03-28 & 2454188 & \\nodata & 0.96 & 2.85 $\\pm$ 0.011 & 6.52 $\\pm$ 0.516 \\\\ XZ Tau & 3531776 & 2004-03-04 & 2453069 & 2 $\\times$ 6 sec & -0.30 & 2.35 $\\pm$ 0.008 & 4.43 $\\pm$ 0.163 \\\\ \\nodata & 15111680 & 2005-09-12 & 2453626 & 8 $\\times$ 6 sec & 0.46 & 2.37 $\\pm$ 0.006 & 1.29 $\\pm$ 0.087 \\\\ \\nodata & 15118336 & 2006-03-09 & 2453804 & \\nodata & 0.36 & 2.42 $\\pm$ 0.007 & 4.16 $\\pm$ 0.090 \\\\ \\nodata & 15118592 & 2006-03-17 & 2453812 & \\nodata & 0.66 & 2.44 $\\pm$ 0.009 & 4.04 $\\pm$ 0.212 \\\\ \\nodata & 16875776 & 2006-03-21 & 2453816 & \\nodata & 0.75 & 2.42 $\\pm$ 0.006 & 3.67 $\\pm$ 0.238 \\\\ \\nodata & 19494144 & 2006-10-17 & 2454026 & \\nodata & -0.32 & 2.50 $\\pm$ 0.007 & 6.12 $\\pm$ 0.094 \\\\ \\nodata & 19493888 & 2006-10-22 & 2454031 & \\nodata & -0.57 & 2.50 $\\pm$ 0.007 & 5.86 $\\pm$ 0.236 \\\\ \\nodata & 19493632 & 2007-03-22 & 2454182 & \\nodata & 0.26 & 2.34 $\\pm$ 0.010 & 3.92 $\\pm$ 0.066 \\\\ \\nodata & 19493376 & 2007-03-28 & 2454188 & \\nodata & 0.10 & 2.34 $\\pm$ 0.008 & 3.97 $\\pm$ 0.046 \\\\ \\enddata  \\tablenotetext{a}{Errors for the color temperature are based off of relative flux uncertainties.} \\tablenotetext{b}{Errors for the flux of the 10~$\\mu$m silicate feature are based off of pointing-induced flux uncertainties.} \\end{deluxetable*} ",
        "conclusions": "The 10~$\\mu$m silicate emission complex observed in the mid-infrared spectra of classical TTS and formed by super-heated dust grains in the upper atmospheres of the innermost regions of their disks can vary on month-long timescales.   Better temporal resolution is required to determine if the variability is periodic.  Our spectral decomposition models were incapable of fitting the silicate feature at every phase of modulation, suggesting that variability may be an explanation for some of the difficulties previous groups have encountered with modeling single-epoch spectra of individual stars.  The simple two-slab two-temperature model suggests an optical depth effect may explain the observed variations.  Since the intervening cooler material may be of similar composition, the variability may not directly impact the disk evolutionary sequence previously constructed for TTS, based on amorphous versus crystalline abundances and dust grain sizes.  However, we note that these models require a number of parameters and the question of degenerate solutions needs to be addressed.  Additional multi-epoch observations of these sources and others are required to determine the fraction of disks with variable silicate emission and provide more variations to model, helping to constrain acceptable parameters.  With the loss of IRS on \\emph{Spitzer}, this project will turn to ground-based observations with less sensitivity, but still the ability to measure the previously observed silicate variability."
    },
    "0910/0910.1384_arXiv.txt": {
        "abstract": "Is the turbulence in cluster-forming regions internally driven by stellar outflows or the consequence of a large-scale turbulent cascade? We address this question by studying the turbulent energy spectrum in NGC~1333. Using synthetic $^{13}$CO maps computed with a snapshot of a supersonic turbulence simulation, we show that the VCS method of Lazarian and Pogosyan provides an accurate estimate of the turbulent energy spectrum. We then apply this method to the $^{13}$CO map of NGC~1333 from the COMPLETE database. We find the turbulent energy spectrum is a power law, $E(k)\\propto k^{-\\beta}$, in the range of scales 0.06~pc $\\le \\ell \\le 1.5$~pc, with slope $\\beta=1.85\\pm 0.04$. The estimated energy injection scale of stellar outflows in NGC~1333 is $\\ell_{\\rm inj}\\approx 0.3$~pc, well resolved by the observations. There is no evidence of the flattening of the energy spectrum above the scale $\\ell_{\\rm inj}$ predicted by outflow-driven simulations and analytical models. The power spectrum of integrated intensity is also a nearly perfect power law in the range of scales 0.16~pc$<\\ell<$7.9~pc, with no feature above $\\ell_{\\rm{inj}}$. We conclude that the observed turbulence in NGC~1333 does not appear to be driven primarily by stellar outflows. ",
        "introduction": "The structure of density and velocity fields in giant molecular clouds can be characterized by extended power laws (e.g. Heyer and Brunt 2004; Padoan et al. 2004, 2006), with scaling exponents consistent with those of numerical simulations of supersonic turbulence (e.g. Padoan et al. 2007; Kritsuk et al. 2007, 2009a,b). The existence of these scaling laws and of one-point statistics of turbulent flows (for example the probability distribution of gas density) is an important assumption in statistical theories of star formation aimed at the prediction of the stellar initial mass function (Padoan and Nordlund 2002, 2004; Hannebelle and Chabrier 2008, 2009) and the star formation rate (Krumholz and McKee 2005; Padoan and Nordlund 2009).  Is this assumption of a universal large-scale turbulent cascade valid in cluster-forming regions, believed to be the birthplace of most stars, or is the turbulence there driven internally by stellar outflows? As an alternative to the idea that stellar clusters are formed in one crossing time (Elmegreen 2000), Tan, Krumholz, and McKee (2006) propose that stars are formed in protocluster clumps over several dynamical times, while the turbulence is driven by stellar outflows. This scenario is investigated analytically by Matzner (2007) and numerically by Nakamura and Li (2007), Carroll et al. (2007), and Wang et al. (2009).  Nakamura and Li (2007) find that the velocity power spectrum in their outflow-driven simulation is very shallow above a characteristic scale of energy injection by outflows, $\\ell_{\\rm{inj}}$. They derive $\\ell_{\\rm inj}\\approx0.3$~pc, assuming their computational domain has a size $L=1.5$~pc and a total mass $m_{\\rm tot}=929$~m$_{\\odot}$, characteristic of cluster-forming regions. This value of $\\ell_{\\rm inj}$ and the flattening of the power spectrum at larger scales is consistent with the analytical results in Matzner (2007). Carroll et al. (2009) confirm the results of Nakamura and Li (2007) at higher numerical resolution, producing velocity power spectra even more clearly peaked at the outflow driving scale. With the values of mass, column density, and outflow momentum per unit mass, $v_{\\rm c}=50$~km/s, adopted by Nakamura and Li (2007), equation (38) in Matzner (2007) gives $\\ell_{\\rm inj}=0.27$~pc. The injection scale is only weakly dependent on $v_{\\rm c}$, $m_{\\rm tot}$, and the star formation rate, and is approximately given by $\\ell_{\\rm inj}\\approx0.079\\,\\Sigma_{\\rm cgs}^{-15/28}$~pc, where $\\Sigma_{\\rm cgs}$ is the column density in g/cm$^2$. Our conclusions do not critically depend on a precise knowledge of $\\ell_{\\rm inj}$, as long as scales just above $\\ell_{\\rm inj}$ are resolved in the observations.  The prototype of outflow-driven cluster-forming regions chosen by Nakamura and Li (2007), Matzner (2007), and Carroll et al. (2009) is NGC~1333 in Perseus. In this Letter, we derive the power spectrum of integrated intensity, $P_I(k)$, and velocity, $P_v(k)$, in NGC~1333, based on the data from the Five College Radio Astronomy Observatory (FCRAO) survey of the Perseus molecular cloud complex (Ridge et al. 2006), publicly available from the COMPLETE website. We show that the scale $\\ell_{\\rm inj}$ is resolved by the observations, and both power spectra, $P_I(k)$ and $P_v(k)$, are consistent with power laws extending to large scale, with the same slope found in supersonic turbulence simulations and in the Perseus complex on larger scale (Padoan et al. 2006). The constant slopes of $P_v(k)$ and $P_I(k)$ above the scale $\\ell_{\\rm inj}$ suggest that outflows may not be the dominant driving mechanism. In a recent paper appeared after the submission of this Letter, Brunt, Heyer, and Mac Low (2009) have reached the same conclusion for NGC~1333 and for other molecular cloud regions using a different method. ",
        "conclusions": "It has been suggested that protocluster clumps actively form stars for several dynamical times, while being supported by the turbulence internally generated by stellar outflows (e.g. Tan, Krumholz, and McKee 2006). This scenario is simulated by Nakamura and Li (2007) and Carroll et al. (2009), with physical parameters characteristic of NGC~1333. They find an energy injection scale, $\\ell_{\\rm inj}\\approx0.3$~pc, above which the turbulent energy spectrum is very shallow, up to the size of 1.5~pc of their computational domain. This value of $\\ell_{\\rm inj}$ and the flattening of the power spectrum above that scale are consistent with the results of the analytical model of Matzner (2007). If the turbulence in cluster-forming regions is primarily driven by outflows, the turbulent energy spectrum should flatten above the scale $\\ell_{\\rm inj}$, up to a scale where the external turbulent cascade dominates again.  We have searched for such a feature in the energy spectrum of NGC~1333, the prototype of outflow-driven regions in the cited theoretical works. To this aim, we have applied the VCS method of Lazarian and Pogosyan (2006) to the $^{13}$CO FCRAO map of NGC~1333, after successfully testing the method on simulated data. We have found the energy spectrum is a power law, $E(k)\\propto k^2P_v(k)\\propto k^{-\\beta}$, with $\\beta=1.85\\pm0.04$, in the range of scales 0.06~pc $\\le \\ell \\le 1.5$~pc, with no evidence of internal driving by outflows at $\\ell_{\\rm inj}=0.3$~pc. The power spectrum of integrated intensity is also a power law, $P_I(k)\\propto k^{-\\beta_I}$, with $\\beta_I=3.19\\pm0.04$, in the range of scales 0.16~pc$<\\ell<$7.9~pc, with no significant features above the predicted injection scale.  These power spectra are consistent with those of large-scale driven simulations of supersonic turbulence, and with those measured on larger scale in Perseus  and other molecular cloud complexes. Although outflows from young stars are present in NGC~1333, the large-scale turbulent cascade appears to be the main energy source. The turbulence in NGC~1333 is either currently driven by significant mass inflow from larger scales, or in the process of being dissipated, until the time when winds, outflows, and ionizing radiation from stars will completely disperse the star-forming gas. \\\\"
    },
    "0910/0910.1517_arXiv.txt": {
        "abstract": "The new era of software signal processing has a large impact on radio astronomy instrumentation. Our design and implementation of a 32 antennae, 33 MHz, dual polarization, fully real-time software backend for the GMRT, using only off-the-shelf components, is an example of this. We have built a correlator and a beamformer, using PCI-based ADC cards and a Linux cluster of 48 nodes with dual gigabit inter-node connectivity for real-time data transfer requirements. The highly optimized compute pipeline uses cache efficient, multi-threaded parallel code, with the aid of vectorized processing. This backend allows flexibility in final time and frequency resolutions, and the ability to implement algorithms for radio frequency interference rejection. Our approach has allowed relatively rapid development of a fairly sophisticated and flexible backend receiver system for the GMRT, which will greatly enhance the productivity of the telescope. In this paper we describe some of the first lights using this software processing pipeline. We believe this is the first instance of such a real-time observatory backend for an intermediate sized array like the GMRT. ",
        "introduction": "\\label{intro} Radio astronomy started with single dishes. The sensitivity and resolution of a single dish radio telescope is limited by its physical aperture area. Due to the larger wavelengths involved, the resolution of even the largest such radio telescope -- the Arecibo dish, which is 300 m in diameter -- is poor compared to its optical counterparts. To overcome this limitation, radio astronomers have evolved the technique of interferometry which synthesizes large apertures using multiple single dishes. The first multi-element radio interferometer for astronomy was built by Sir Martin Ryle \\cite{Ryle}. An interferometer measures the cross-correlation between the signals for a given pair of antennae. The instantaneous value of this cross-correlation (also called ``visibility'') estimates one Fourier component of the two dimensional brightness distribution in the sky plane.  The frequency of the measured Fourier component depends on the vector distance between the two antennae, projected onto a plane normal to the direction of the source (also called ``baseline\"). For an array of $N$ antennae, at any given time, there will be $N*(N-1)/2$ such baselines and corresponding visibility measurements. Due to the rotation of the Earth, the baseline vectors change with time, allowing many more Fourier components to be measured.  By taking a 2-D Fourier Transform of the measured visibilities (after appropriate calibration), one can obtain the sky brightness distribution \\cite{Thompson}, with a resolution that corresponds to an aperture size equal to the largest separation in the array. Some of the major aperture synthesis interferometer arrays in operation today are, the Very Large Array (VLA) of NRAO, the Westerbork Synthesis Radio Telescope (WSRT) of ASTRON, the Australia Telescope Compact Array (ATCA) of ATNF, the Giant Metrewave Radio Telescope (GMRT) of NCRA.  The GMRT consists of an array of 30 antennae, each of 45 m diameter, spread over a region of 25 km diameter, and operating at 5 different wave bands from 150 MHz to 1450 MHz \\cite{Swarup}.  Amongst the multi-element earth rotation aperture synthesis telescopes operating at meter wavelengths, the GMRT has the largest collecting area. The GMRT can also be configured in array mode, where it acts as a single dish by adding the signals from individual dishes \\cite{Gupta}. This mode of operation is used for studying compact objects like pulsars, which are effectively point sources even for the largest interferometric baselines. The maximum instantaneous operating bandwidth at any frequency band is 33 MHz. Each antenna provides signals in two orthogonal polarizations, which are processed through a heterodyne receiver chain and brought to the central receiver building, where they are converted to baseband signals and fed to the digital backend consisting of correlator and pulsar receiver. The existing GMRT hardware backend uses  Application Specific Integrated Circuit based design approach, which is fairly efficient and optimized for its prime scientific goals, but has limited flexibility. It also suffers from some limitations due to quantization effects in the hardware pipeline.  The pulsar receiver, which uses Digital Signal Processing chips, is somewhat more flexible in its configuration and capabilities.  With the growing impact of the GMRT, there is a strong need felt by the user community for a flexible and more powerful backend which can support new regimes of observations. Some examples of these are : (i) spectral line observations requiring higher frequency resolution over larger bandwidths  (ii) pulsar searches and surveys for transient sources requiring higher time resolution across a wide field of view, including the capability of forming multiple phased array beams within the primary field of view.  Furthermore, like most low frequency radio telescopes, the GMRT is also affected by local radio frequency interference (RFI) -- both narrow spectral lines and broadband, impulsive signals. To produce data with high dynamic range and low noise, it is essential to detect and filter out such RFI signals, at various stages in the processing pipeline.  The existing GMRT hardware backend is not designed to address such issues. A software based backend can be an attractive option to overcome these limitations.  A software backend is a real-time or off-line processing pipeline that runs on a supercomputer or cluster of computers.  The rapid growth in general purpose computing power during the last few years had made it possible to compete with the processing speeds of dedicated hardware. Besides this, fast data transport links between computers are now possible due to the availability of high speed gigabit networking. Storage media have also gone through a revolution in the last few years, in terms of capacity and throughput. All these recent advancements have made it attractive to attempt designs of software backends for radio astronomy.  This is aided by the fact that the computing required for the processing of signals from a multi-element interferometric array is very well suited for implementation on multi-processor computers. Use of commercial off-the-shelf (COTS) components and software processing blocks significantly reduces the several years of development time required for hardware backends.  It also makes the upgrades much easier, by replacing with faster computers without spending effort in rebuilding the processing blocks again. The algorithms implemented in software are much more flexible; new features can be added easily, and new regions of parameter space can be explored very conveniently.  Design of correlators using off-the-shelf general purpose computers is not an entirely 21st century approach. Very long baseline interferometric (VLBI) observations, where the distances between the antennae are much greater than in conventional interferometers, were started with the recording of raw voltage samples from each antenna onto magnetic media. The recorded data from all the antennae in the VLBI network were correlated off-line using general purpose computers. The first such software correlator was implemented on IBM mainframes (system 360/50) at the NRAO by the NRAO-Cornell VLBI group around 1967 \\cite{Bare}. The input data were from a Mark I tape recording system, albeit with a rather modest operating bandwidth of 360 kHz \\cite{Kellermann}. Over the next few decades, the increase in complexity of the VLBI systems -- due to increase in bandwidths and number of antennae $-$ resulted in a switch towards specialized hardware correlators to process the data, e.g. Very Long Baseline Array (VLBA) of NRAO, Joint Institute for VLBI in Europe (JIVE), the Canadian S2 system.  This was the era where hardware/firmware options were found to provide realistic solutions \\cite{Carlson}.  However, Moore's law was slowly catching up with these requirements, and by the late 1990s, the availability of computing power (including distributed parallel processing) and high speed, high density data storage capabilities renewed the interest of the radio astronomy community in software processing. In 1990, a gated cross-correlator for Mark II VLBI was devised by Petit et al. \\cite{Petit}. The DiFX VLBI correlator \\cite{Deller} designed by the Swinburne University of Technology, which extensively employs parallel processing techniques, is the most recent milestone in this regard. However, all these projects are mostly for off-line processing of recorded data.  Phillips et al. \\cite{Phillips} demonstrated real-time e-VLBI for an array of four telescopes from the Australian Long Baseline Array, with 16 MHz of observing  bandwidth, using the DiFX correlator.  However, the work described here to develop a software based backend for the GMRT is probably the first implementation of a fully {\\it real-time} software pipeline as a regular observatory backend for a medium sized array. The upcoming BlueGene-L based central processing unit for the LOFAR telescope \\cite{Romein} will be another such fully real-time backend. The IBM cell processor based e-VLBI correlator is also aimed for real-time data transport and software based processing \\cite{Wagner}.  The recently developed GMRT software backend (GSB), built using mainly COTS components, is a fully real-time backend that supports all the features of the existing hardware backend of the GMRT. In addition, it provides substantially enhanced spatial and temporal resolutions.  Furthermore, it supports a baseband recording mode where raw voltage signals from all the antennae can be recorded to disk for specialized off-line processing to maximize the science return.  In this paper, we describe our design and implementation of this 32 antennae, 33 MHz, dual polarization, fully real-time software backend for the GMRT. ",
        "conclusions": "\\label{sec:19}  We have described our design and implementation of a 33 MHz bandwidth, 32 station, dual polarization, fully {\\it real-time} software backend system for the GMRT. Our approach has allowed relatively rapid development of a fairly sophisticated and flexible backend receiver system, which will greatly enhance the productivity of the GMRT.  We have successfully validated the backend and it has now been released for regular use. We have also demonstrated some of the versatile features of the new backend, and described its capabilities for RFI mitigation and other enhanced features.  We believe this is the first instance of a software based ${\\it real-time}$ backend for an intermediate sized array like the GMRT.  Our approach holds promise for future developments for bigger radio telescopes and wider bandwidths."
    },
    "0910/0910.4935_arXiv.txt": {
        "abstract": "{ Under the assumption that the variations of parameters of nature and the current acceleration of the universe are related and governed by the evolution of a single scalar field, we show how information can be obtained on the nature of dark energy from observational detection of (or constraints on) cosmological variations of the fine structure constant and the proton-to-electron mass ratio. We also comment on the current observational status, and on the prospects for improvements with future spectrographs such as ESPRESSO and CODEX. ",
        "introduction": "We propose to probe the dynamics of the equation of state of matter and energy in the universe by using observations of cosmologically varying fundamental parameters. This requires a number of assumptions: first, that there is clear evidence for dark energy that could be attributed to a rolling scalar field or quintessence; second, that there is a cosmological variation of the fine structure constant of magnitude $\\Delta\\alpha/\\alpha \\lesssim 10^{-5}$, as suggested by Keck/HIRES high resolution quasar absorption spectra (\\cite{Murphy:2003mi,Murphy:2003hw}, but see also \\cite{Srianand:2007qk}); and thirdly, that any variation of $\\alpha$ arises from the evolution of the quintessence field (\\cite{Dvali:2001dd,Chiba:2001er}).  Concerning the first assumption, there is currently little observational evidence for dark energy to be more than a bare cosmological constant. However, if dark energy indeed results from an evolving scalar field, it could be expected to couple to other forms of matter and lead to variations of masses and couplings (implying the second and third assumptions) unless some unknown symmetry principle explicitly forbids these couplings.  We take the coupling between the scalar field and electromagnetism to be ${\\cal L}_{\\phi F} = - \\frac{1}{4} B_F(\\phi) F_{\\mu\\nu}F^{\\mu\\nu}$ where the gauge kinetic function $B_F(\\phi)$ is linear, $ B_F(\\phi) = 1- \\zeta \\kappa (\\phi-\\phi_0)$ (and $\\kappa^2=8\\pi G$). This can be seen as the first term of a Taylor expansion, and should be a good approximation if the field is slowly varying at low redshift. Then, the evolution of alpha is given by \\begin{equation} \\frac{\\Delta \\alpha}{\\alpha} \\equiv \\frac{\\alpha-\\alpha_0}{\\alpha_0} = \\zeta \\kappa (\\phi-\\phi_0) \\,. \\end{equation}  We can also consider the variation of $\\mu\\equiv m_p/m_e$. In grand unified theories we expect a correlation between the variation of $\\alpha$ and $\\mu$ (\\cite{Calmet:2001nu}) given by $\\Delta \\mu/\\mu = R\\, \\Delta \\alpha/\\alpha$, where $R$ is a model-dependent numerical factor arising from correlated variations of $\\Lambda_{QCD}$, the Yukawa couplings, the vacuum expectation value of the Higgs field and $\\alpha$ itself. Under simple assumptions we obtain $R \\sim -20$, which is in severe tension with observations that indicate a nontrivial variation of $\\alpha$ at high redshift but null variation of $\\mu$ (\\cite{King:2008ud,Thompson}) with equal or better precision. This simple exercise illustrates the potential of cosmological observations of quasar absorption lines and variation of fundamental parameters in discriminating particle physics models. ",
        "conclusions": "We have shown that under simple assumptions we can determine the nature of dark energy, not by {\\it fitting}\\/ the parameters of a scalar potential to cosmological data, but by performing the inverse procedure, that consists in using quasar data to {\\it reconstruct}\\/ the potential. We have seen that the evolution of the equation of state parameter can in principle be found, subject of course to the precision of future data. These observations have profound implications as the simple knowledge of the sign of $w'(z)$ can help us to favour or discard freezing models of quintessence (increasing $w(z)$ with increasing redshift), thawing models (decreasing $w(z)$) and k-essence models (typically also decreasing $w(z)$).  This type of reconstruction directly probes the scalar field dynamics and may be carried out, with current data, to redshifts beyond $z = 4$, far higher than the limiting value $z = 1.7$ of SnIa searches. Future astrophysical techniques may extend this to even higher redshifts (\\cite{Levshakov,Kozlov}). Moreover, the observations can be done from the ground and are consequently cheaper than satellite-based observations.  Here, we have presented the reconstruction procedure for a minimally coupled scalar field, but other models with non-canonical kinetic terms, couplings to matter or multiple fields might have further interesting phenomenological properties and therefore lead to alternative approaches."
    },
    "0910/0910.5722_arXiv.txt": {
        "abstract": "Because cosmic superstrings generically form junctions and gauge theoretic strings typically do not, junctions may provide a signature to distinguish between cosmic superstrings and gauge theoretic cosmic strings. In cosmic microwave background anisotropy maps, cosmic strings lead to distinctive line discontinuities. String junctions lead to junctions in these line discontinuities. In turn, edge detection algorithms such as the Canny algorithm can be used to search for signatures of strings in anisotropy maps. We apply the Canny algorithm to simulated maps which contain the effects of cosmic strings with and without string junctions. The Canny algorithm produces edge maps. To distinguish between edge maps from string simulations with and without junctions, we examine the density distribution of edges and pixels crossed by edges. We find that in string simulations without Gaussian noise (such as produced by the dominant inflationary fluctuations) our analysis of the output data from the Canny algorithm can clearly distinguish between simulations with and without string junctions. In the presence of Gaussian noise at the level expected from the current bounds on the contribution of cosmic strings to the total power spectrum of density fluctuations, the distinction between models with and without junctions is more difficult. However, by carefully analyzing the data the models can still be differentiated. ",
        "introduction": "In recent years there has been a revival in interest in finding signatures of cosmic strings. One of the main reasons is that a large class of string inflationary models predict \\cite{Sarangi:2002yt} a copious production of cosmic superstrings \\cite{Witten:1985fp} - cosmic strings consisting of extended fundamental strings of cosmological scale - at the end of the period of inflation. Under certain conditions \\cite{Copeland:2003bj}, these cosmic superstrings are stable on cosmological time scales.  Since gauge theories also predict the presence of cosmic strings \\cite{Kibble} (see \\cite{VilShell,HK,RHBrev} for reviews), it is an important endeavor to distinguish between generic cosmic strings and cosmic superstrings. As we shall demonstrate in this note, the plethora of small scale Cosmic Microwave Background (CMB) experiments, such as the ground based Atacama Cosmology Telescope (ACT) \\cite{ACT} and the South Pole Telescope (SPT) \\cite{SPT} or the Planck satellite experiment \\cite{Planck}, will provide the observational means to make this distinction.  Both generic cosmic strings and cosmic superstrings leave imprints in the CMB.  The key signature - the ``Kaiser-Stebbins (KS) effect'' \\cite{Kaiser} - consists of line discontinuities in the temperature map formed from a combination of gravitational lensing and the Doppler effect: photons from the last scattering surface streaming by either side of a moving cosmic string will be observed to have a temperature which differs by a small amount proportional to the string tension. The signature is manifest in position space - in Fourier space the phase information which is crucial to obtain a line discontinuity is lost. Therefore, to find evidence for the line discontinuities, new data analysis techniques are required which work directly in position space. One such technique is the Canny edge detection algorithm \\cite{Canny1,Canny2}. As shown in \\cite{Amsel,Stewart,Danos:2008fq}, the Canny algorithm can be used to search for the line discontinuities produced by strings. In fact, it was shown that when applied to data with angular resolution comparable to that which will be obtained from the South Pole Telescope, the non-detection of KS lines will allow an improvement in the upper bound on the cosmic string tension by almost an order of magnitude compared to existing bounds.  One essential distinction between strings in simple gauge theory models and cosmic superstrings is that the latter come in different varieties \\footnote{For sufficiently complicated gauge theory models, it is also possible to obtain different types of strings \\cite{Bucher}. Fat strings arising in theories with light moduli fields also can give rise to junctions \\cite{Wells} - in this case due to mergings of strings with low winding number into strings of higher winding number. For simple gauge theory strings, strings with higher winding number are unstable to the fragmentation into strings of winding number 1, and hence these junctions to not arise.}. There can be fundamental strings (F-strings), one space-dimensional branes (D-strings), and bound states of the above. These flavors of cosmic superstrings can form junctions, points where three strings meet in the shape of a Y \\cite{Copeland:2003bj,Leblond:2004uc,Jackson:2004zg,Hanany:2005bc, Hashimoto:2005hi,Copeland:2006eh,Copeland:2006if,Copeland:2005cy,Tye:2005fn}.   In the present paper we find that the Canny algorithm can be modified in order to allow a differentiation between string maps with and without string junctions. Since string junctions are generically predicted in models of brane inflation with cosmic superstrings but are not present for cosmic strings in simple gauge theory models, our work points towards a way of finding observational evidence which would favor cosmic superstrings over simple gauge theory strings.  To identify the line discontinuities predicted by cosmic strings, good angular resolution of a CMB experiment is crucial. It is for this reason that our simulations work with maps similar in angular extent and angular resolution which will be obtained by the ACT and SPT experiments, those with ideal angular resolution.  Our implementation of the Canny algorithm \\cite{Danos:2008fq} finds line discontinuities produced by the KS effect and produces a map of edges. In contrast to previous work \\cite{limits} on constraining the cosmic string tension from CMB observations which focused on the angular power spectrum of CMB maps in which the KS signature is washed out, our work searches specifically for the line discontinuities predicted by the KS effect.  Our numerical work is based on simulations of the predicted CMB anisotropies in small patches of the sky (side length $10^{o}$) with angular resolution of $1.5^{'}$. We simulate maps predicted in pure string models (i.e. no ``Gaussian noise'') with and without junctions, and also maps in which the ``Gaussian noise'' due to inflationary perturbations dominates the total amplitude of the power spectrum, and cosmic strings contribute a sub-dominant fraction consistent with the current upper bounds \\cite{limits}. The current upper bound on the contribution of strings to the angular power spectrum of CMB anisotropies corresponds to a value of the string tension $\\mu$ which is $G \\mu < 2 \\times 10^{-7}$ in dimensionless units ($G$ is Newton's gravitational constant).  Our implementation of the Canny algorithm consists of a set of routines which generate a map of edges in the sky corresponding to gradients in CMB maps which are in the range expected from the KS effect \\cite{Danos:2008fq}, and statistically analyze the resulting edge maps, with the goal of distinguishing between simulated CMB maps of pure Gaussian inflationary perturbations and maps combining simulated cosmic strings with Gaussian perturbations. In our studies \\cite{Danos:2008fq} we showed that this distinction could be made up to a three sigma confidence level for strings with tensions as low as $2\\times 10^{-8}$ for maps without the removal of point source noise (unsmoothed) and as low as $9\\times 10^{-8}$ for maps with point source noise removed (smoothed).  Here we extend the use of the Canny algorithm to distinguish simulated maps of strings with and without junctions. Based on the edge maps which the Canny algorithm produces, we count the number of edges present in groups of pixels, thus determining the density of edges.  We then use the distribution of densities to differentiate between maps with and without junctions. Specifically, we compare the shapes of the histograms (number of occurrences versus number of edges) in the two cases, using the ``combined Fisher probability method''.  The simulations contain various free parameters. Firstly, there is the cosmic string tension $\\mu$. Secondly, there is the number $N$ of string segments per Hubble volume per Hubble time during the string scaling phase. The fraction of string segments which involve a junction is another choice which must be made. In order to ascertain that a deviation in the shape of an edge histogram from that in a theory with pure Gaussian noise is a consequence of strings with junctions, we must allow the parameters in the simulations with and without junctions to be different. For example, it is not sufficient to show that the histogram in a string simulation with junctions for $N = 10$ is different from that of a string simulation without junctions for the same value of $N$. We must also check that the change in the shape resulting from the addition of junctions cannot be masked by a change in the value of $N$. We have carefully studied this point and find that, with appropriately chosen number of boxes sampled in the observed window, we are able to differentiate string simulations with and without junctions. Without Gaussian noise the difference in the shape of the histogram is manifest. However, with Gaussian noise consistent with the cosmic strings contributing not more than they are allowed to by the limits of \\cite{limits} we find that the difference is no longer manifest. However, with sufficient care, we are still able to make the statistical distinction.  The outline of this paper is as follows: we first describe our simulated microwave sky patches resulting from models with cosmic strings junctions (Section II).  In Section III we review the application of the Canny algorithm and present the edge maps which result from our simulations. We present the density distribution analysis in Section IV. Finally, in Section V we conclude with a discussion of the results. ",
        "conclusions": "We have applied the Canny algorithm to simulated maps of CMB anisotropies induced by models with cosmic stings either containing or not containing string junctions. We proposed a statistic with which it should be possible to distinguish between the maps with and without junctions. This provides a method with which one can distinguish between the predictions of simple gauge field theory models with cosmic strings (which typically do not admit string junctions), and models such as those giving rise to cosmic superstrings, which are characterized by the presence of junctions.  We first showed that our statistic is able to clearly differentiate between string maps with and without junctions in the absence of Gaussian noise. However, since we know that a cosmic string model without a dominant Gaussian contribution to the spectrum of fluctuations is not consistent with the latest data on the angular power spectrum of CMB anisotropies, we need to consider correctly normalized sky maps which contain strings with a tension sufficiently low such that its contribution to the power spectrum does not exceed the current limits \\cite{limits}, with the bulk of the power coming from Gaussian noise. In this case it is more difficult to differentiate between maps where the strings have junction and where they do not.  Nevertheless, making use of our detailed statistics, we found a clear distinction in the shape (peak and width) of the density distribution (number of boxes sampled with varying numbers of edges or pixels in the edges) for the different scaling solutions.  For example, $N=1$ strings per Hubble volume had a large number of boxes with just a few edges, as would be expected, whereas $N=10$ strings per Hubble volume looked more Gaussian, as expected.  We found a statistical difference between the curves for junctions and the curves of different scaling solutions without junctions.  Future work includes optimizing the Canny algorithm, and searching for the lowest value of the string tension for which maps with and without string junctions can be differentiated. It would also be very interesting to algorithms to ``real'' string simulations (those obtained via a numerical evolution of the Nambu-Goto equations) rather than just relying on the simple toy models for the distribution of strings which we have used. Maps of CMB anisotropies produced by strings without junctions coming from ``real'' simulations are available \\cite{Fraisse}. However, to our knowledge there are no corresponding results for networks with string junctions."
    },
    "0910/0910.0448.txt": {
        "abstract": "{}{}{}{}{} % 5 {} token are mandatory  \\abstract {It is debated whether the Milky Way bulge has the characteristics of a classical bulge sooner than those of a pseudobulge. Detailed abundance studies of bulge stars is a key to investigate the origin, history, and classification of the bulge. These studies can give constraints on the star-formation history, initial mass-function, and trace differences in stellar populations. Not many such studies have been made due to the large distance and large variable visual extinction along the line-of-sight towards the bulge. Therefore, near-IR investigations are to be preferred.} {The aim is to add to the discussion on the origin of the bulge and to study detailed abundances determined from near-IR spectra for bulge giants already investigated with optical spectra, the latter also providing the stellar parameters which are very significant for the results of the present study. Especially, the important CNO elements are better determined in the near-IR. Oxygen and other $\\alpha$ elements are important for the investigation of the star-formation history. The C and N abundances are important for determining the evolutionary stage of the giants but also the origin of C  in the bulge.} {High-resolution, near-infrared spectra in the H band are recorded using the CRIRES spectrometer on the {\\it Very Large Telescope}. The CNO abundances can all be determined from the numerous molecular lines in the wavelength range observed. Abundances of the $\\alpha$ elements Si, S, and Ti are also determined from the near-IR spectra. } {[O/Fe],  [Si/Fe] and [S/Fe] are enhanced up to metallicities of at least [Fe/H]=$-0.3$, after which they decline. This suggests that the Milky Way bulge experienced a rapid and early star-formation history like that of a classical bulge. However, a similarity between the bulge trend and the trend of the local thick disk seems present. Such a similarity could suggest that the bulge has a pseudobulge origin. The C and N abundances suggest that our giants are first-ascent red-giants or clump stars, suggesting that the measured oxygen abundances are those the stars were born with. Our [C/Fe] trend does not show any increase with [Fe/H] which could have been expected if W-R stars have contributed substantially to the C abundances. No \"cosmic scatter\" can be traced around our observed abundance trends; the scatter found is expected, given the observational uncertainties.} %. Accurate fundamental parameters are important, but difficult to achieve. } % results heading (mandatory) % {The CNO abundances are simultaneously needed since their abundances influence each other strongly through molecular equilibria. {} ",
        "introduction": "One great unsolved question in cosmology is how galaxies formed and acquired their stellar populations \\citep[see for instance][]{renzini}. Two scenarios have been traditionally entertained for the formation of bulges, with one  leading to ``classical\" bulges, whereby they form from merger-driven starbursts, and another leading to ``pseudobulges\" that would result from the ``secular\", dynamical evolution of disks \\citep{kormendy}. In the classical bulge scenario most stars originate in a short phase of star formation when the universe was only a few Gyr old, and bulge formation may precede that of the disk, that could be acquired later. Instead, in the so-called ``pseudobulges\" stars form in the disk, over an extended period of time, and the bulge results from the secular evolution of the disk driven by the development of a bar. Thus, one expects classical bulges to be made almost exclusively of old stars, while stars in pseudobulges would have an age spread comparable to the Hubble time.  The bulges of early-type spirals (Sa and Sb) are generally considered to be ``classical\", while pseudobulges are thought to inhabit preferentially late-type spirals (Sc and Sd). The Sbc Milky Way (MW) galaxy sits morphologically on the borderline, hence it is no surprise that the origin of its bulge is currently a matter of debate. What is especially puzzling with the MW bulge is that it looks dynamically ``pseudo'' (with its boxy, peanut-shaped bar, Kormendy \\& Kennicutt 2004), while its stellar content is just what one expects for a ``classical'' bulge. Indeed, deep color-magnitude diagrams (CMD) of the MW bulge show no detectable trace of stellar populations younger than $\\sim10$~Gyr (Ortolani et al. 1995; Zoccali et al. 2003)\\nocite{ortolani,zoccali:2003}, a conclusion for which compelling evidence has been recently provided by the deep CMD obtained with HST for a proper-motion selected sample of bulge stars \\citep{clarkson}.  Such a  sharp distinction between these two scenarios of bulge formation has been called into question recently: observations of disk galaxies at redshift $\\sim 2$ (lookback time $\\sim 10$~Gyr) have indeed revealed that such early disks have radically different physical properties compared to local disks of similar mass. The early disks are characterized by much higher velocity dispersion and gas fraction, and harbor massive, highly star-forming clumps \\citep{genzel:06,forster}. Thus, dynamical instabilities in such early disks would occur with much shorter timescales (few $10^8$ years) compared to local disks, with a secular, but fast evolution of such disks resulting in the early formation of a bulge \\citep{genzel:08}.  In parallel to observations, theories have also envisaged an early production of bulges via the migration and central coalescence of gas-rich clumps in high-redshift disks \\citep[e.g.,][]{immeli,elmegreen,carollo,bournaud}.  Determinations of detailed chemical compositions are key data for studies of the origin and evolution of stellar populations, since they carry characteristic signatures of the objects that enrich the interstellar gas. Abundance ratios are sensitive to the time-scales of star formation, to the initial mass function (IMF), and may disclose relations between different stellar groups, since different elements are synthesized by different processes and stars. For the bulge, a high ratio of $\\alpha$-element abundances relative to Fe is observed, suggesting that the star-formation period was early and very short \\citep{lecureur,fulbright:07} and that the bulge formed more rapidly than the thin, and perhaps even the thick galactic disk  \\citep{mcwilliam:08}.  With stellar abundance ratios carrying such a {\\it genetic} information on the origin of different stellar populations, \\citet{zoccali} compared the [O/Fe] ratios of bulge stars with those of thin disk and thick disk stars from \\citet{bensby:04}. Finding bulge ratios systematically higher than disk ones, Zoccali et al. argued against the bulge having been appreciably contaminated by the migration of thick disk stars analogous to those in the solar neighborhood.  A similar consideration was put forward by \\citet{lecureur}, based on the much higher values of their [Na/Fe], [Mg/Fe] and [Al/Fe] ratios in bulge stars compared to those in thin and thick disk stars from \\citet{bensby:05}, \\citet{bensby:kol}, and \\citet{reddy}. However, differences in abundance ratios may be also due to differences in systematic errors when the abundance analysis is done by different groups, using different data and methods. For example, cool giants with crowded spectra are prone to many systematic uncertainties \\citep[see, for instance,][]{santos:09} and comparing abundances derived from metal-rich, cool giants to those of solar-type dwarfs could be uncertain due to relative systematics.   In this context, Mel\\'endez et al. (2008) carried on an homogeneous abundance analysis of bulge, thick disk and thin disk giant stars, confirming the [O/Fe] trend found by Zoccali et al. (2006) for bulge stars, but finding [O/Fe] ratios for thick disk stars much higher than those derived by Bensby et al. (2004). Thus, in the study by Mel\\'endez et al. the high thick-disk [O/Fe]-ratios appear to be indistinguishable from those of the bulge. This similarity of [O/Fe] ratios between bulge and thick disk stars weakens the conclusion of Zoccali et al. (2006) about the genetic difference between the bulge and the thick disk, although it remains to be seen whether this similarity holds also for other [$\\alpha$/Fe] ratios, that appear so different in Lecureur et al. (2007).  If thick disk and bulge giants follow the same [O/Fe] vs [Fe/H] relation the inference would be that both have formed \"rapidly\". The properties of z$\\sim2$ disks, with their high velocity dispersion and high star formation rate, suggest that what we see there is thick-disk formation.   Determination of abundances for a large sample of red giant stars and planetary nebulae \\citep[cf.][]{chiappini:09} in various bulge fields as well as in the inner galactic disk  will obviously provide a most powerful method to constrain the chemical evolution and models of the bulge \\citep{matteucci:99,silk}. A way to achieve this is by high-resolution, near-IR spectroscopy of red-giant stars. In the IR, the obscuration in the direction of the bulge is considerably reduced, offering us  the opportunity to go to heavily reddened regions. %We need to explore various galactic longitudes to discriminate between signatures of a pseudo and a classical bulge. Furthermore,  near-IR spectra suffer much less from line blending than spectra at optical wavelengths, which makes it possible to safely trace the continuum and avoid abundance criteria marred with blending lines, so important in abundance analysis. Moreover, only the IR offers, even within a small wavelength range, all indicators necessary to accurately determine the CNO abundances by the simultaneous observations of many clean CO, CN and OH lines. Here, we present the first data from our VLT/CRIRES programme in which we systematically study stellar abundance ratios from different parts of the Galactic bulge. In particular, we have aimed at the key elements C, N, and O, but also include some $\\alpha$ elements. %, but we have also determined %abundances of several other important %elements:  Fe, Ti, Ni, Si, and S. ",
        "conclusions": "Abundance determinations for stars are known to be plagued by systematic errors that may be difficult to estimate. In order to discuss the properties of different stellar populations, homogeneous differential spectroscopic studies, and detailed comparisons of results of different studies, are significant. In the present study we have tried to follow this route, and find a satisfactory agreement with results obtained in the optical, as well as IR, when a common temperature scale is used for the stars. With our high-resolution IR spectroscopy, we have explored the CNO abundances, as well as the abundances of Si, S, Ti and Fe, for 11 bulge giants. We have found enhanced [O/Fe], [Si/Fe], and [S/Fe] values with increasing [Fe/H] up to approximately [Fe/H]$\\sim-0.3$, after which these abundance ratios relative to Fe decrease. This suggests an early and rapid star formation in the bulge. Our investigation is not devised to make a detailed comparison with  thick disk stars and determine the relationship between these two populations; such a study should be made differentially to minimise systematic uncertainties.  Our abundance trends are, however, consistent with a similarity between these populations as was found in the differential study  by \\citet{melendez:2008}. %Alves-Brito et al. (2009, in prep.). Such a similarity suggests that the picture of an isolated classical bulge may be oversimplified. Inner-disk stars, at smaller galactocentric distances, should be explored in order to deepen the understanding of a possible physical connection between the bulge and the thick disk. %If there is a similarity, then the bulge star-formation rate and initial-mass function would not need to be  particular  compared to the local %thick disk \\citep{melendez:2008}. This would then imply a possibility for the galactic bulge to have been dynamically formed from inner disk %stars through transient  disk instabilities, as in the formation of a pseudo-bulge.  From our C and N abundances we conclude that our stars are first-ascent red-giants or clump stars, suggesting that their oxygen abundances are not affected by CNO cycling. Furthermore, we find that  there is no significant increase in the carbon abundances for high metallicities, which could have been expected if  W-R stars were to explain the large decline in [O/Fe] vs. metallicity.  We have demonstrated that  for the same stars several different determinations of the stellar parameters from optical spectra produce significantly different results, implying important systematic uncertainties. Attempts to reduce these should be made. Note also that  \\citet{chiappini:09} compare, among others, the oxygen abundances derived from planetary nebulae (PNe) and giants in the bulge and find that the abundances determined from  giant star spectra are systematically higher by 0.3 dex. They conclude that this discrepancy may be caused by systematic uncertainties in either the PNe or giant star abundance determinations, or both.  To fully clarify the situation of the origin and evolution of the galactic bulge, additional near-IR abundance surveys of  elements (especially more $\\alpha$ elements) are needed. Most earlier investigations have been restricted to Baade's window. Different regions of the bulge are now being explored, and further systematic such work is needed.  %Attempts to trace the results of a mixing-in of other populations than the dominant old one should be made. %A statistically significant sample of stars will then be needed to enable detection of a possible %minor young secular population of stars.  A proper age classification of the stellar populations in the Milky Way %bulge will then be possible.   %%%%%%%%%%%%  %"
    },
    "0910/0910.4160.txt": {
        "abstract": "An intriguing possibility for TeV scale physics is the existence of neutral long lived particles (LOLIPs) that subsequently decay into SM states. Such particles are many cases indistinguishable from missing transverse energy (MET) at colliders. We propose new methods to search for these particles using neutrino telescopes. We study their detection prospects, assuming production either at the LHC or through dark matter (DM) annihilations in the Sun and the Earth. We find that the sensitivity for LOLIPs produced at the LHC is limited by luminosity and detection energy thresholds. On the other hand, in the case of DM annihilation into LOLIPs, the sensitivity of neutrino telescopes is promising and may extend beyond the reach of upcoming direct detection experiments. In the context of low scale hidden sectors weakly coupled to the SM, such indirect searches allow to probe couplings as small as $10^{-15}$. ",
        "introduction": "\\label{sec:introduction} Current and future experiments are beginning to probe the TeV scale. This energy scale incorporates the mechanism for Electroweak symmetry breaking (EWSB) and quite possibly the solution to the dark matter (DM) mystery. \u00ca While many theoretical possibilities for TeV scale physics have been explored in great detail, it is still quite possible that something truly unexpected will be discovered. \u00ca \u00caThe reason for this is because no experimental evidence thus far that favors a particular model proposed so far. The possibility for exotic experimental signatures from beyond the SM physics has been recently explored in more detail in an attempt to ensure that nothing is missed at the LHC. \u00caHidden Valleys~\\cite{Strassler:2006im}, Quirks~\\cite{Kang:2008ea} and other theories have demonstrated the limitations of current experimental searches at colliders by exploring the possibility that new physics might not be triggered on, or easily reconstructed at the LHC.  One intriguing possibility for new physics are particles which have very long lifetimes. \u00caIf these particles are charged under the SM, they are typically known as charged massive particles (CHAMPs).  These have striking experimental signatures, but are straightforward to search for. \u00caIf the LOLIPs are neutral, experimental searches are much more difficult and strongly depend on their lifetime. \u00caIn this paper we will explore new ways to search for neutral LOLIP decays that allow one to explore lifetimes that were previously beyond the experimental reach.  Neutral LOLIPs are quite generic in models of physics beyond the SM. \u00caFor instance, additional gauge sectors or new particles that couple only very weakly to the SM can easily have macroscopic lifetimes, $\\tau \\gg \\mathcal{O}(10^{-6})\\,\\mathrm{s}$. \u00caThe possibility for such lifetimes has already been discussed in many specific contexts before. \u00caFor instance, in theories of \u00cagauge mediated supersymmetry breaking (GMSB), the decay length of the next to lightest super-partner (NLSP) into the gravitino is macroscopic when the primordial supersymmetry (SUSY) breaking scale is $\\sqrt{F}\\gtrsim 1000\\,\\mathrm{TeV}$. \u00caAnother scenario, that of a low scale hidden sector, has received considerable attention during the last year. \u00ca Indeed, indirect clues from astrophysics have motivated studies of these ``unconventional\" models at the TeV scale. \u00caThe recent PAMELA~\\cite{pamela} and Fermi~\\cite{fermi} measurements imply that if the experimental signals are a consequence of DM annihilations, DM must have non-standard properties~\\cite{us}, suggesting the possible existence of an additional sector, weakly coupled to the SM~\\cite{nima}. \u00caIf this is the case, the low lying states in the hidden sector may be long lived and decay to the SM through some small coupling, $\\epsilon$.  Recently several ways have already been proposed to directly search for neutral LOLIP decays when the decay length is not much longer than the size of the detector used. \u00ca These fall under the general categories of using high energy colliders~\\cite{ArkaniHamed:2008qp, Han:2007ae, feng}, low energy $e^+e^-$ colliders~\\cite{Essig:2009nc,reece,Aubert:2009pw}, fixed target experiments~\\cite{Bjorken:2009mm,reece, pospelovFT}, and rare meson decays~\\cite{reece}. \u00caWhile many new ideas for searching for neutral LOLIPs have been put forth, the lifetimes probed are relatively short and correspondingly the coupling $\\epsilon$ is typically not smaller than $10^{-7}$ at most.  An even more intriguing and less studied possibility is the case when the decay length of a neutral LOLIP is much greater than the size of the detector used to produce and study it. \u00caFor instance if neutral LOLIPs are produced at the LHC and have an effective decay length significantly larger than the size of ATLAS or CMS, they will simply escape the detector and be recorded as MET. \u00caFrom this perspective one could not distinguish these particles from stable particles.  Experimental ways to differentiate between these two possibilities are very limited. \u00caTwo indirect methods have been suggested for probing a hidden sector that has particles with very long lifetimes: (i) looking at the effects on the evolution in stars, for instance supernova cooling~\\cite{Bjorken:2009mm} and (ii) measurements of certain gamma ray features in models related to DM~\\cite{Ruderman:2009ta,Ruderman:2009tj}. \u00caThese methods are very model dependent and in this paper we will present more generic methods that allow one to search for LOLIPs with decay lengths much greater than $\\mathcal{O}(\\mathrm{km})$ .  There are three ingredients which dictate the possible reach of a search for LOLIPs. \u00caThe first is the production mechanism and its rate. \u00caThis mechanism will also determine the momentum distribution of the LOLIPs, which may influence the capability to detect this particle. \u00caThe second ingredient is the lifetime of the particle, which determines the probability of the LOLIPs to decay inside the detector. \u00ca \u00caThe physical quantity of interest is the effective lifetime of the LOLIPs, $\\tau_{eff}=\\gamma \\tau$, which depends on the typical momenta of the produced particle. \u00ca \u00caFinally, the final states into which the LOLIP decays, will determine the types of detectors that are most applicable and how these particles can be searched for.  In this paper we will focus on two mechanisms for production. \u00caThe first, and most straightforward perhaps, is to produce the particle at a collider. \u00caIf the LOLIP weakly couples to the SM but has a low mass, then in principle high luminosity low energy experiments can produce it~\\cite{reece}. \u00caAlternatively, the LOLIP could have a larger mass, as is the case for instance, of a bino NLSP in gauge mediation. Perhaps more interestingly, the long lived state might be in another sector that is only effectively accessed by producing heavier states such as a Z' or supersymmetric particles, which subsequently decay into this new sector. \u00caIn this case luminosity is still important, but higher energy is critical for producing these new states. \u00caThus the most effective way for producing \u00cathese particles is with the LHC. \u00ca  The second method, which turns out to be more promising, is producing LOLIPs from DM annihilation. \u00ca For this production mechanism to be useful, one has to focus on regions where the DM density is high, such that a large enough flux is generated. \u00caThe most interesting regions that are potentially useful are the Sun and Earth. \u00ca\u00caGiven the association with DM, there are a number of already existing constraints on this scenario, coming from Fermi~\\cite{fermi}, Milagro~\\cite{Atkins:2004qr}, and SuperKamiokande~\\cite{upgoingmuons} that \u00caexclude part of the parameter space.  For either production mechanism, we show that using large volume neutrino telescopes is the most powerful method for increaseing the reach to longer lifetimes and weaker couplings . \u00caAs we will demonstrate, these experiments can probe hidden sectors with $\\epsilon$ as small as $10^{-15}$, and lifetimes comparable to an Astronomical Unit, in the case of production through DM. \u00caWe also demonstrate that detecting the decay of LOLIPs produced at the LHC is very difficult even with the improved upcoming experiments. \u00caHowever we suggest new methods that may be improved upon, and in addition put forth the idea of new detectors dedicated to these searches that will be investigated further in~\\cite{workinprogress}.  The rest of the paper is structured as follows. \u00caIn Section~\\ref{sec:models} we begin by discussing specific examples of models that contain neutral LOLIPs. \u00caWe give the parametric dependence of their decay lengths, which is necessary to translate into specific model parameters the reaches that we present as a function of lifetime. \u00caIn Section~\\ref{sec:experiments} we give the details for the neutrino telescope experiments that we will use to bound neutral LOLIP decays and to search for them. \u00caWe include the details of how neutral LOLIP decays differ from the standard events looked at by these experiments and quantify this difference for certain cases. \u00ca In Section~\\ref{sec:lhc} we discuss potential methods for discovering neutral LOLIPs produced at the LHC. \u00caIn Section~\\ref{sec:sunandearth} we discuss the bounds and the discovery potential for neutral LOLIP decays coming from DM annihilation in the Earth and Sun. \u00caIn that Section, we also discuss the bounds from the Fermi and Milagro experiments in the case that LOLIPs are produced from DM. \u00caFinally in Section~\\ref{sec:conclusions} we discuss future avenues for LOLIP searches and the ramifications of the bounds and reaches \u00cawe calculated for individual models. ",
        "conclusions": "\\label{sec:conclusions}   In this paper we have studied the prospects of detecting neutral long lived particles (LOLIPs) with decay lengths ranging between $1\\unit{km}$ and $10^{15}\\unit{km}$. \u00ca \u00caSuch long lived states may arise in several scenarios such as GMSB models and hidden sectors weakly coupled to the SM. \u00ca Previous studies focused on significantly shorter lifetimes, keeping many theoretical models out of the experimental reach. \u00caWe stress that the added value of the approach discussed in this paper, is that it not only improves on the sensitivity to low scale weakly coupled sectors, but also it allows for detection of hidden sectors with mass scales significantly larger than $1\\unit{GeV}$, such as the model (III-d) of Section~\\ref{sec:models}.  To obtain sensitivity to such long lifetimes, we studied neutrino telescopes employing their large scale detectors. \u00caFor the production mechanism, we considered the LHC and DM annihilations into LOLIPs in the Sun or the Earth. \u00ca \u00caAttempting to directly probe LOLIPs produced at the LHC and detected at Antares or KM3NeT proved not feasible due to the low rates. \u00caFuture improvements on the luminosity at the LHC and/or on the energy threshold at KM3NeT may render this possibility viable. \u00caConversely, designing small scale detectors closer to the LHC may allow one detect LOLIPs which show up as MET at the LHC. \u00caWe postpone the study of this possibility to future work~\\cite{workinprogress}. On the other hand, indirect detection of LOLIPs from the Earth and Sun seems promising, depending on the final states into which the LOLIPs decay. \u00ca In particular, if LOLIPs predominantly decays into di-muons, a significant \u00caregion of the parameter space previously unexplored can be probed by IceCube. \u00caConversely, the prospects of detecting final states which shower inside the detector, such as electrons or photons, strongly depend on the capabilities of reducing the atmospheric neutrino background. \u00ca   Finally, in relation to DM annihilating into LOLIPs, \u00cawe have compared the capabilities of the neutrino telescope techniques, to those of direct detection. We found that for models which accommodate long lived states, neutrino telescopes can detect annihilation fluxes from the Sun corresponding \u00cato nucleon cross-sections beyond the reach of upcoming direct detection experiments.  Note Added:  While this work was in writing, related work appeared~\\cite{Batell:2009zp}.  {\\bf Acknowledgments.} We thank N.~Arkani-Hamed, J.~Ruderman, P.~Schuster, M.~Strassler, N.~Toro, J.~Uscinski and I.~Yavin for useful conversations.   PM, MP, and TV would like to thank the Galileo Galilei Institute for Theoretical Physics in Florence, Italy where part of this work was completed.  MP would like to thank the Aspen Center for Physics where part of this work was completed. PM and TV are supported in part by DOE grant DE- FG02-90ER40542. MP is supported in part by NSF grant PH0503584."
    },
    "0910/0910.1792_arXiv.txt": {
        "abstract": "Since the Newtonian gravitation is largely used to model with success the structures of the universe, such as galaxies and clusters of galaxies, for example, a way to probe and constrain alternative theories, in the weak field limit, is to apply them to model the structures of the universe. We then modified the well known \\textbf{Gadget-2} code to probe alternative theories of gravitation through galactic dynamics. In particular, we modified the \\textbf{Gadget-2} code to probe alternatives theories whose weak field limits have a Yukawa-like gravitational potential. As a first application of this modified \\textbf{Gadget-2} code we simulate the evolution of elliptical galaxies. These simulations show that galactic dynamics can be used to constrain the parameters associated with alternative theories of gravitation. ",
        "introduction": "\\label{intro} Studying Cosmology is a way to probe the physical origin, the present, and destiny of the Universe. The present Cosmological scenario claims that the Universe has the following composition: $\\sim 5 \\%$ by barionic matter, $\\sim 25 \\%$ by dark matter, and $\\sim 70 \\%$  of a component that works like an antigravity term in the Einstein's Equations: the dark energy.  Dark matter can be explained by physically acceptable arguments (e.g., WIMPs or modified theory of gravitation). Forthcoming experiments to be made in the Large Hardron Collider (LHC) will certainly help, in the near future, to give some answers concerning this questions (whether the WIMPs be detected or not).  On the other hand, the dark energy composition is harder to explain, because it can not be thought as due to particles, and its interpretation through a quantum vacuum gives a very bad answer: the energy scale is more than a hundred orders of magnitude greater than the observed value (about this  discrepancy, see, e.g., \\cite{cpt1992}). It is very interesting to note that scalar field explanations - such as in some quintessence models - do not solve the problem at all, because other questions arise concerning the very origin of quintessence and its nature.  In this puzzling scenario, some physical theories have been created to explain alternatively the dark energy problem and the dark matter nature. These alternative approaches are not nonsense, because until now, dark matter can only be detected indirectly \\textit{via} observations of galactic or extragalactic gravitational interactions, unless LHC changes this situation. These theories are based on different hypothesis (e.g., scalar-tensor theories of gravity,  massive gravitons, etc.) and they do not have in the weak field limit the Newtonian gravitational potential (see, e.g., Refs. \\cite{moffat0,moffat1,moffat2,moffat3,moffat4,piazza03,rodriguez-meza05,sign05,araujo2007}).  In some cases, the gravitational potential is Yukawa-like (hereafter Yukawian gravitational potential, YGP). Considering the most simple form of this potential, if we have a point mass $m$, namely  \\begin{equation} \\label{ygp} \\phi=-\\frac{Gm}{r}e^{-r/\\lambda}, \\end{equation} where $r$ is the distance from the point mass, and $\\lambda$ is a characteristic length, that means, in some theories, the Compton wavelength of the exchange particle of mass $m_{\\rm g}$, a  massive boson called graviton.  It is worth noting that, in almost all works concerning the YGP, the investigations have been made under analytical or numerical approaches. Graviton masses are estimated over many ranges, in different scales of observations, based on different scenarios, from planetary to extragalactic scales. Besides, statical models are commonly used, without taking into account any additional investigations concerning, for example, secular dynamics of the systems.  In this way, we propose numerical simulations of triaxial systems, to probe the YGP, as given by Eq. \\ref{ygp}; and investigate how different the galaxies would appear if the YGP were more ``realistic'' than the Newtonian potential at large scales. In other words, if YPG can not keep the King Sphere in dynamical equilibrium, for example, it must be ruled out, because it would destroy the Hernquist spheres and the exponential disks as well, consequently galaxies could not exist in this scenario.  The paper is organized as follows: In Section 2, we present the code and the galactic model used, in Section 3, the simulations and results and, finally, in Section 4, we discuss the results and show the perspectives. ",
        "conclusions": "In this work, we have studied triaxial systems to probe the YGP and to constrain the yukawian $\\lambda$ parameter. We showed here that, if YGP were reliable, we should have $\\lambda > 100$ kpc. This value is larger then that from solar system's constraints and it must be considered a good estimative, since with such a value the simulated galaxies remains ``alive\" for billions of years and look like their observational counterparts. It is worth stressing that triaxial systems easily constrain the $\\lambda$ parameter, because we only need to study the radial density profile, particles's positions and the core's velocity dispersion. Although this procedure is somewhat simplified, it give us sufficient information about the status of the model's structure and its observational counterparts. In particular, models presenting the destruction of a galaxy, such that depicted in the first frame of Fig. \\ref{snaps_king}, or presenting bizarre morphologies mean that the parameter under consideration is unreliable.  This work has some important characteristics. We have designed a method to probe alternative theories of gravitation in the \\textit{non relativistic r\\'egime}, simulating ``alive\" galaxies submitted to the investigated potential. This method reveals itself as a test to constraint the parameters of the probed theory. Concerning the alternative theories of gravitation, we have developed a pioneer simulation method probe in the sense that previous works have studied statical models (e.g., using galactic scenario, de Araujo \\& Miranda studied disk galaxies under YGP, but using analytical arguments), while our galaxies behave like ``alive\" systems, because are composed by thousands of gravitating particles. In this way, this must be considered a reliable and strong test, due to the fact that N-Body systems are very sensitive to the physics used in the simulation, and their observational counterparts are required to give us the best parameters of the potential under investigation. Although chaos and complex phenomena appears in the N-Body systems, it is important to bear in mind that these systems are well understood by Galactic Dynamics (see, e.g., \\cite{bt2008}).  Another interesting test would be making disk galaxies under the YGP. This test would be a more sophisticated one, due to the fact that late-type systems have complex substructures that are expected to appear in these simulations, like spiral arms, bars and rings. These features can be used to probe the $\\lambda$ values with more precision. In other papers to appear elsewhere \\cite{brandaoearaujo2009b}, \\cite{brandaoearaujo2009c} we will make these simulations under YGP and other potentials."
    },
    "0910/0910.4381_arXiv.txt": {
        "abstract": "The discovery of very high energy (VHE) gamma-ray emitting X-ray binaries has triggered an intense effort to better understand the particle acceleration, absorption, and emission mechanisms in compact binary systems, which provide variable conditions along eccentric orbits. Despite this, the nature of some of these systems, and of the accelerated particles producing the VHE emission, is unclear. To answer some of these open questions, we conducted a multiwavelength campaign of the VHE gamma-ray emitting X-ray binary \\lsi\\ including the MAGIC telescope, \\xmm, and \\swift\\ during 60\\% of an orbit in 2007 September. We detect a simultaneous outburst at X-ray and VHE bands, with the peak at phase 0.62 and a similar shape at both wavelengths. A linear fit to the simultaneous X-ray/VHE pairs obtained during the outburst yields a correlation coefficient of $r=0.97$, while a linear fit to all simultaneous pairs provides $r=0.81$. Since a variable absorption of the VHE emission towards the observer is not expected for the data reported here, the correlation found indicates a simultaneity in the emission processes. Assuming that they are dominated by a single particle population, either hadronic or leptonic, the X-ray/VHE flux ratio favors leptonic models. This fact, together with the detected photon indices, suggests that in \\lsi\\ the X-rays are the result of synchrotron radiation of the same electrons that produce VHE emission as a result of inverse Compton scattering of stellar photons. ",
        "introduction": "\\label{introduction}  \\lsi\\ is one of the few X-ray binaries that have been detected in very high energy (VHE) gamma rays (see, e.g., \\citealt{paredes08} for a recent review). It is a high-mass X-ray binary system located at a distance of 2.0$\\pm$0.2~kpc \\citep{frail91}. The system contains a rapidly rotating early type B0\\,Ve star with a persistent equatorial decretion disk and mass loss, and a compact object with a mass between 1 and 4~$M_\\odot$ orbiting it every $\\sim$26.5~d in an eccentric orbit (see \\citealt{casares05}; \\citealt{grundstrom07}; \\citealt{aragona09}, and references therein). \\lsi\\ was classified as a microquasar based on structures detected from five to several tens of milliarcseconds with the European VLBI Network (EVN) and MERLIN (\\citealt{massi04}, and references therein), although analysis of later EVN and MERLIN data sets does not reveal the presence of such structures \\citep{dhawan06,albert08_lsi2006_mw}. Very Long Baseline Array (VLBA) images with a resolution of 1~mas (2~AU at the source distance) obtained during a full orbital cycle show an elongated structure that rotates as a function of the orbital phase \\citep{dhawan06}. Later VLBA images show repeating structures at the same orbital phases, reinforcing the idea that the milliarcsecond structure depends on the orbital phase \\citep{albert08_lsi2006_mw}. This may be consistent with a model based on the interaction between the relativistic wind of a young non-accreting pulsar and the wind/decretion disk of the stellar companion \\citep{dubus06a}, as occurs in PSR~B1259$-$63 (but confining the pulsar wind may be problematic; see \\citealt{bogovalov08}).  \\lsi\\ shows periodic non-thermal radio outbursts on average every $P_{\\rm orb}$=26.4960$\\pm$0.0028~d, with the peak of the radio emission shifting progressively from phase 0.45 to 0.95, using $T_0$=JD~2,443,366.775, with a modulation period of 1667$\\pm$8~d \\citep{gregory02}. According to the most precise orbital parameters, the periastron takes place at phase 0.275 and the eccentricity of the orbit is $0.537\\pm0.034$ \\citep{aragona09}.  \\lsi\\ has been observed several times in X-rays (see \\citealt{smith09} and references therein). It generally displays quasi-periodic X-ray outbursts, with the maximum occurring between orbital phases 0.4 and 0.8 \\citep{goldoni95,taylor96,paredes97,harrison00,esposito07}, although the lack of a sensitive long-term monitoring has prevented to search for a super-orbital modulation. The source also shows short-term variability on timescales of hours \\citep{sidoli06}.  At VHE gamma rays, \\lsi\\ has been clearly detected several times both by MAGIC \\citep{albert06_lsi_science,albert08_lsi2006_mw,albert09_lsi2006_period} and VERITAS \\citep{acciari08_lsi,acciari09_lsi}. The source also displays VHE gamma-ray periodicity, with minima taking place near periastron, where only upper limits were found in MAGIC observations, and maxima occurring on average at phase 0.6--0.7, although the source has also shown a second peak at phase 0.84 in one of the cycles \\citep{albert09_lsi2006_period}. There are indications of correlated X-ray/VHE emission, based on non-simultaneous data taken more than 6~hr apart \\citep{albert09_lsi2006_period}, and 1~d apart \\citep{albert08_lsi2006_mw}.  The lack of a systematic behavior from cycle to cycle at X-ray and VHE bands, and the occurrence of short-term variability in the X-ray flux, has hampered the definitive detection of an X-ray/VHE correlation from the comparison of non-simultaneous data. We therefore conducted a multiwavelength campaign in 2007 September, covering the epoch of maximum VHE emission. Here, we report the first simultaneous VHE and X-ray observations of \\lsi\\ obtained with the MAGIC Cherenkov telescope and the \\xmm\\ and \\swift\\ X-ray satellites that show correlated emission in both energy bands. We also briefly comment on radio and optical spectroscopic observations obtained during the campaign. ",
        "conclusions": "\\label{discussion}  We have discovered an X-ray/VHE gamma-ray correlation in the gamma-ray binary \\lsi\\ based on simultaneous multiwavelength data obtained with MAGIC, \\xmm, and \\swift\\ during a single orbital cycle. The correlation is mainly due to the very similar trends of the detected flux during the outburst around orbital phase 0.62, while the uncertainties prevent to be sure about the existence of the correlation outside the outburst. Given the variability in the X-ray flux (up to 25\\% on hour scales), it is necessary to program simultaneous observations to perform correlation studies in \\lsi. This conclusion has also been reached recently by the VERITAS Collaboration, when reporting the lack of X-ray/VHE correlation based on contemporaneous data with a VHE sampling that is not dense enough \\citep{acciari09_lsi}. Although a similar X-ray/VHE correlation has been obtained for the VHE gamma-ray emitting X-ray binary \\object{LS~5039}, this result was based on non-simultaneous data acquired years apart \\citep{aharonian05_ls5039_science,aharonian06_ls5039_period,takahashi09}.  The VHE emission within a binary system can suffer photon--photon absorption via pair creation, mainly with the stellar optical/ultraviolet photons. In \\lsi, this absorption is only expected to be significant toward the observer for $E>300$~GeV just before periastron \\citep{dubus06b,bednarek06,sierpowska09}, a phase range not explored here. In addition, the quoted X-ray fluxes are already de-absorbed (and the hydrogen column densities and the associated errors are low). Overall, there are no absorption effects to be considered. Therefore, the X-ray/VHE correlation we have found for \\lsi\\ cannot be an artifact due to variable absorption toward the source. This indicates that the emission processes at both wavelengths occur at the same time and are probably the result of a single physical mechanism. In this context, it is reasonable to assume that the X-ray and VHE emissions are produced by a single particle population.  It is interesting to note that the MAGIC spectrum in the 0.6--0.7 phase range yields $\\sim11\\times10^{-12}$~erg~cm$^{-2}$~s$^{-1}$ for $E>300$~GeV, while the X-ray flux is $\\sim19\\times10^{-12}$~erg~cm$^{-2}$~s$^{-1}$. Therefore, the total X-ray flux is approximately twice the VHE flux. However, if we subtract an apparent baseline X-ray flux of $10\\times10^{-12}$~erg~cm$^{-2}$~s$^{-1}$, the resulting X-ray flux is similar to the total VHE flux in the phase range 0.6--0.7. If the radiation mechanisms are dominated by a single particle population, the X-ray/VHE correlation and the smaller/similar VHE fluxes favor leptonic models. In hadronic models, the X-ray emitting $e^\\pm$ and the VHE photons would come from the same protons (for reasonable values of the magnetic field), and the luminosity of the $e^\\pm$ radiation should be $\\la 1/2$ that of VHE gamma-rays (see Figure~5 of \\citealt{kelner06} for reasonable proton energy distributions) unlike it is observed. In addition, the inverse Compton (IC) cooling channel is less efficient than the synchrotron channel to produce the detected X-ray emission for reasonable values of the magnetic field (see \\citealt{takahashi09} for a similar discussion for \\object{LS~5039}). This clearly suggests that the X rays are the result of synchrotron radiation of the same electrons that produce VHE emission as a result of IC scattering of optical/ultraviolet stellar photons.  The observed photon indices of the simultaneous X-ray and VHE spectra are consistent with one population of electrons following a power-law energy distribution with index $\\sim$2.1. These electrons would produce X-rays via synchrotron and VHE photons via IC with an interaction angle $\\la \\pi/2$. We note that an electron index of $\\sim$2.1 is too hard if synchrotron cooling dominates in the X-ray range, since it implies an injected electron index of 1.1. On the other hand, dominant adiabatic cooling implies an injection index of 2.1, a more reasonable value (see, for example, the discussion in \\citealt{takahashi09}). Although small changes in the VHE spectrum would be expected due to variations in the IC interaction angle or the electron index (as seen in X-rays), at present they are not detectable due to the large uncertainties of the VHE photon index.  Finally, we note that contemporaneous radio light curves obtained with RATAN, VLBA images, and H$\\alpha$ spectroscopy are consistent with previous results \\citep{gregory02,dhawan06,zamanov99}. Details on these observations will be reported elsewhere. Therefore, the X-ray/VHE correlation occurred when the source was showing a standard behavior in both its outflow (radio) and decretion disk (H$\\alpha$ line)."
    },
    "0910/0910.3935_arXiv.txt": {
        "abstract": "Mid-infrared (IR) observations of polycyclic aromatic hydrocarbons (PAHs) and molecular hydrogen emission are a potentially powerful tool to derive physical properties of dense environments irradiated by intense UV fields. We present new, spatially resolved, \\emph{Spitzer} mid-IR spectroscopy of the high UV-field and dense photodissocation region (PDR) around Monoceros R2, the closest ultracompact \\hII region, revealing the spatial structure of ionized gas, PAHs and H$_2$ emissions. Using a PDR model and PAH emission feature fitting algorithm, we build a comprehensive picture of the physical conditions prevailing in the region. We show that the combination of the measurement of PAH ionization fraction and of the ratio between the H$_2$ 0-0 S(3) and S(2) line intensities, respectively at 9.7 and 12.3 $\\mu$m, allows to derive the fundamental parameters driving the PDR: temperature, density and UV radiation field when they fall in the ranges $T = 250-1500\\,$K, $n_H=10^4-10^6$cm$^{-3}$, $G_0=10^3-10^5$ respectively. These mid-IR spectral tracers thus provide a tool to probe the similar but unresolved UV-illuminated surface of protoplanetary disks or the nuclei of starburst galaxies. ",
        "introduction": "\\label{int}   Dense ($>10^4$ cm$^{-3}$) and high UV field ($>10^4$ times the standard interstellar radiation field in units of the Habing field, written $G_0$ hereafter) photodissociation regions rule the energy balance, and thus evolution, of some of the most fundamental astrophysical objects such as protoplanetary disks, planetary nebulae and starburst galaxies. One of the best tools to probe this UV illuminated matter is spectroscopy in the mid-infrared (5-15 $\\mu$m, hereafter mid-IR), because it provides information on both the gas and small dust grain properties (polycyclic aromatic hydrocarbons and very small grains, hereafter PAHs and VSGs) . Ideally, one would like to spatially resolve the emission of these components in order to study their variations as the UV field is attenuated. Unfortunately, such observations are very hard to achieve (for the moment) on one hand because large ground-based telescopes, providing arcsecond angular resolution, are restrained in wavelength coverage, while on the other hand, space borne telescopes have diameters that are usually too small to resolve the sources. Because it is the closest ultra compact \\hII region at a distance of 850 pc, Monoceros R2 (Mon R2, see \\citealt{woo89} and \\citealt{how94}) constitutes one of the rare exceptions where one can resolve the PDR between the \\hII region and molecular cloud. In this Letter, we present and analyze the mid-IR PAH and molecular hydrogen emissions in Mon R2 based on spatially resolved \\emph{Spitzer} spectral mapping.   \\begin{figure*} \\begin{center} \\epsscale{1.0} \\plotone{MonR2_CS+NE+PAH+H2_3.eps} \\vspace{-0.0cm} \\caption{Overview of the Mon R2 region seen with \\emph{Spitzer}. \\emph{Left:} Observed IRS spectra for PDR 1, 2 and 3 (continuous line) and fit by the model (diamonds, see Sect.\\ref{mim}). \\emph{Middle:}  The blue contours represent the intensity of the \\neII line (0.02 to 0.1 erg s$^{-1}$ cm$^{-2}$sr$^{-1}$ in linear steps), the red contours the intensity of the  PAH 11.3 $\\mu$m band (1.8 to 3.4 $\\times 10^{-2}$ erg s$^{-1}$ cm$^{-2}$sr$^{-1}$ in linear steps) and the green contours the intensity of the H$_2$ 0-0 emission in the S(3) rotational line at 9.7 $\\mu$m (1.5 to 4.5 $\\times 10^{-4}$ erg s$^{-1}$ cm$^{-2}$sr$^{-1}$ in linear steps). The background map shows the CS $J$=5-4 emission presented in \\citet{cho00}. Low emission is in black.  Numbers indicate the positions of the selected PDRs. \\emph{Right:} Cut in the map along the yellow box, showing the stratified emission of the different lines. \\label{Mfit} } \\end{center} \\end{figure*} ",
        "conclusions": "The main results we have presented in this Letter are: (\\emph{i}) Mon R2 is a unique template of high-UV/high-density PDRs that is spatially resolved, similar to the Orion bar (where the UV field is slightly lower), (\\emph{ii}) our observations are consistent with PDR models which predict that in high-UV and dense PDRs, the excitation of pure rotational H$_2$ lines is collisional and depends on gas temperature, (\\emph{iii}) for this reason, the spatial distributions of PAHs and warm H$_2$ are different, contrary to what is seen in cool PDRs, (\\emph{iv}) however, because PAH emission is present in the extended region illuminated by the radiation, their ionization fraction can be used to probe the intensity of the UV radiation field, even where H$_2$ emission peaks. Thus, the mid-IR spectrum of dense and highly irradiated PDRs appears as an efficient probe of the physical conditions in these environments. The ratio between the molecular H$_2$ 0-0 S(3)/S(2) line intensities allows to directly compute the density and temperature of the gas, while the measurement of the ionization fraction of PAHs allows to derive the intensity of the radiation field. We show that the derived parameters for the entire Mon R2 region are consistent with those found for the spatially resolved PDRs 2 and 3. This is likely because the dense PDR1 occupies only a small part of the whole region, and therefore dilution effects imply that the spatially averaged spectrum is dominated by emission from lower density PDRs. Thus, we suggest that this spectral methodology could be useful to derive the dominant physical conditions in PDRs that are not spatially resolved in the mid-IR. In particular, the PDRs at the surface of protoplanetary disks around Herbig Ae$/$Be stars \\citep{ber09} or in the inner rim of T-Tauri disks \\citep{agu08} are expected to be the targets where such an analysis can be applied. The mid-IR spectrum of starburst galaxies is probably dominated by the emission of PDRs having similar conditions as those described in this paper (see \\citealt{fue08} for the case of M82). Thus, the present methodology could be a useful tool to derive global properties of galaxies, in connection with massive star formation activity, using forthcoming James Webb  and SPICA space telescopes that will observe these emission features at low and high redshifts."
    },
    "0910/0910.3332_arXiv.txt": {
        "abstract": "Moderately close binaries are a special class of targets for planet searches. From a theoretical standpoint, their hospitality to giant planets is uncertain and debated. From an observational standpoint, many of these systems present technical difficulties for precise radial-velocity measurements and classical Doppler surveys avoid them accordingly. In spite of these adverse factors, present data support the idea that giant planets residing in binary and hierarchical systems provide unique observational constraints on the processes of planet formation and evolution. The interest and the importance of including various types of binary stars in extrasolar planet studies have thus grown over time and significant efforts have recently been put into: (i) searching for stellar companions to the known planet-host stars using direct imaging, and (ii) extending Doppler planet searches to spectroscopic and moderately close visual binaries. In this contribution we review the observational progresses made over the past years to detect and study extrasolar planets in binary systems, putting special emphasis on the two developments mentioned above. ",
        "introduction": "\\label{intro}  Nearby G-K dwarfs, which are first-choice targets for many planet search programs, are more often found in binary or multiple systems than in isolation \\citep[e.g.][]{Duquennoy91,Eggenberger04b}. This observation, coupled with the growing evidence that many young binaries possess circumstellar or circumbinary disks susceptible of sustaining planet formation \\citep[e.g.][]{Monin07}, raises the question of the existence of planets in star systems of different types.  From a dynamical standpoint, planets residing in binary systems can be found in three different configurations: (1) in circumstellar orbits (i.e. orbiting the primary or the secondary star, also called S-type orbits); (2) in circumbinary orbits (i.e. circling both stars, P-type orbits); or (3) around the L4 or L5 Lagrange point in systems with a very small mass ratio (Trojan planets, L-type orbits). However, the presence of an additional stellar companion may threaten either the formation or the long-term stability of planetary systems, imposing further restrictions. For instance, a stellar companion within $\\sim$100 AU will likely affect -- and possibly inhibit -- the formation of circumstellar giant planets (e.g. \\citeauthor{Nelson00}~\\citeyear{Nelson00}; \\citeauthor{Mayer05}~\\citeyear{Mayer05}; \\citeauthor{Thebault06}~\\citeyear{Thebault06}; see also the contributions by Kley, Marzari, and Th\\'ebault), while a more distant but highly inclined companion can influence the evolution of these planets on secular timescales \\citep[e.g.][]{Innanen97,Takeda05}. These considerations raise two fundamental questions: What types of binary systems do actually host planets in S/P/L-type orbits?, and Are such planets common or rare?  Most of the information we presently have about planets in binaries comes from ``classical'' Doppler planet searches which target nearby G-K dwarfs. This observational material is highly incomplete with respect to the closest binaries, however, because it is difficult to extract precise radial velocities when the two components simultaneously contribute to the recorded flux \\citep{EggenbergerUdry07}. To avoid light contamination at the spectrograph entrance, Doppler surveys exclude from their target samples most -- but usually not all -- double stars closer than $2$-$6^{\\prime\\prime}$ \\citep[e.g.][]{Udry00,Perrier03,Marcy05,Jones06}, systems which we will call moderately close binaries. This strategy implies that classical Doppler programs provide little information about the existence of planets in spectroscopic and visual binaries $\\lesssim$200~AU. To probe the occurrence of planets in these moderately close systems, new methods and alternative detection techniques have been actively developed over the past years.  This chapter is partly a review on the observational efforts to detect and characterize extrasolar planets in binary systems, and partly a description of our own contribution to this research field. In Sect.~\\ref{obs} we present a brief overview of the methods that are being used to detect extrasolar planets in binary and multiple systems. In Sect.~\\ref{classical_doppler} we summarize the information gathered on planets in binaries through classical Doppler planet searches. We then describe our recent efforts to investigate the impact of stellar duplicity on the occurrence and properties of giant planets. This work follows two complementary approaches: searching for stellar companions to the known planet-host stars using direct imaging (Sect.~\\ref{imaging}), and extending Doppler planet searches to spectroscopic binaries (Sect.~\\ref{pib}). We conclude in Sect.~\\ref{conclusion} with a summary of the main results and future perspectives. ",
        "conclusions": "\\label{conclusion}  Over the past seven years, binaries have become increasingly interesting targets for planet searches. From the observational perspective, Doppler surveys have shown that at least 17\\% of the known planetary systems are associated with one or more stellar companions \\citep[e.g.][]{EggenbergerUdry07} and that circumstellar giant planets exist in binaries as close as 20 AU \\citep{Queloz00,Hatzes03,Zucker04,Correia08}. From the theoretical perspective, simulations indicate that the presence of a stellar companion within $\\sim$100~AU likely affects the formation and evolution of circumstellar giant planets \\citep{Kley00,Nelson00,Mayer05,Boss06,Thebault06}, leaving potential imprints in the occurrence and properties of these objects. Circumstellar planets residing in binaries $\\lesssim$100~AU thus provide unique testing grounds for the models of planet formation and evolution.  Imaging and literature surveys searching for stellar companions to the known planet-bearing stars have been very successful, revealing many new binary and multiple planet-host systems \\citep[e.g.][]{Patience02,Raghavan06,Chauvin06,Desidera07,Eggenberger07b,Mugrauer07}. Thanks to its well-defined control sample, our NACO survey provides the first piece of evidence that circumstellar giant planets are intrinsically less frequent in binaries $\\lesssim$100~AU than around single stars \\citep{Eggenberger08}. Future analyses based on larger samples will constrain the dependence of this frequency on binary separation over the range $\\sim$35-250 AU.  The discoveries from classical Doppler planet searches and from imaging surveys indicate that the Kozai mechanism plays a role in shaping the high end of the eccentricity distribution of extrasolar planets \\citep{Wu03,Takeda05,Tamuz08,Moutou09}. Additionally, the distinctive properties of the short-period planets residing in binary or hierarchical systems \\citep{Zucker02,Eggenberger04,Desidera07,Mugrauer07} suggest that some of these planets owe their current orbital configuration to Kozai migration \\citep{Takeda06,Fabrycky07,Wu07}. The data set on planets in binaries thus provides growing evidence that distant stellar companions commonly affect the orbital evolution of planetary systems on secular timescales.  Since 2002 significant efforts have been put into extending planet searches to (moderately) close binaries using different techniques, including Doppler spectroscopy, phase-referenced interferometry, eclipse or pulse timing, transit photometry, gravitational microlensing, and adaptive optics imaging. Although these efforts have not produced many planet discoveries yet, they will yield very important results in the coming years. Upon completion, the ongoing Doppler and interferometric surveys will provide the first measures of the true frequency of circumstellar planets in binaries $\\lesssim$50 AU. When combined with the results from imaging programs, these new data will give some information as to whether circumstellar giant planets commonly form in binaries $\\lesssim$100~AU. At the same time, the present and future Doppler surveys targeting SB2s will quantify the frequency of massive circumbinary planets.  In the next few years, several additional observing facilities will open up new possibilities for planet searches in/around (moderately) close binaries. Thanks to their high photometric precision and continuous monitoring of rich stellar fields, the recently launched CoRoT\\footnote{http://smsc.cnes.fr/COROT/} and Kepler\\footnote{http://kepler.nasa.gov} missions will initiate large-scale planet searches around eclipsing binaries (\\citeauthor{Doyle04}~\\citeyear{Doyle04}; \\citeauthor{Ofir09}~\\citeyear{Ofir09}; see also the contribution by Sybilski). Alternatively, new interferometric facilities like PRIMA at the Very Large Telescope Interferometer \\citep{Delplancke08} or the ASTRA upgrade of the Keck Interferometer \\citep{Pott08} will allow to extend astrometric planet searches to significantly wider and fainter binaries. The upcoming generation of high-contrast adaptive optics instruments such as HiCIAO at Subaru \\citep{Tamura06}, GPI at Gemini South \\citep{Macintosh08}, SPHERE at the Very Large Telescope \\citep{Beuzit08} and PALM-3000/Project-1640 at Palomar \\citep{Bouchez08} will also likely be used to carry out systematic searches for giant planets in some types of moderately close binaries. Gathering together the results from all these different surveys we will then have pretty good answers to the two questions mentioned at the beginning of this chapter: What types of binaries do host planets in S/P/L-type orbits?, and Are such planets common or rare?  As surveys progress and diversify, the conviction that planets are common objects in the universe continually strengthen. In addition to the encouraging observational results obtained so far on circumstellar giant planets, theoretical studies support the existence of low-mass planets in many binary systems (e.g. \\citeauthor{Thebault06}~\\citeyear{Thebault06}; \\citeauthor{Haghighipour07}~\\citeyear{Haghighipour07}; \\citeauthor{Quintana07}~\\citeyear{Quintana07}; see also the contribution by Marzari). Similarly, circumbinary planets are expected to be quite common \\citep[e.g.][]{Quintana06,Pierens07,Scholl07,Pierens08a}. Observational programs targeting (moderately) close binaries thus promise additional exciting results to come. Since a full understanding of planet formation must address the issue of circumstellar and circumbinary planets, these programs are an essential part of today's research on extrasolar planets."
    },
    "0910/0910.2703_arXiv.txt": {
        "abstract": "The accretion-induced collapse (AIC) of a white dwarf (WD) may lead to the formation of a protoneutron star and a collapse-driven supernova explosion. This process represents a path alternative to thermonuclear disruption of accreting white dwarfs in Type Ia supernovae. In the AIC scenario, the supernova explosion energy is expected to be small and the resulting transient short-lived, making it hard to detect by electromagnetic means alone. Neutrino and gravitational-wave (GW) observations may provide crucial information necessary to reveal a potential AIC. Motivated by the need for systematic predictions of the GW signature of AIC, we present results from an extensive set of general-relativistic AIC simulations using a microphysical finite-temperature equation of state and an approximate treatment of deleptonization during collapse. Investigating a set of 114 progenitor models in axisymmetric rotational equilibrium, with a wide range of rotational configurations, temperatures and central densities, and resulting white dwarf masses, we extend previous Newtonian studies and find that the GW signal has a generic shape akin to what is known as a ``Type~III'' signal in the literature.  Despite this reduction to a single type of waveform, we show that the emitted GWs carry information that can be used to constrain the progenitor and the postbounce rotation. We discuss the detectability of the emitted GWs, showing that the signal-to-noise ratio for current or next-generation interferometer detectors could be high enough to detect such events in our Galaxy. Furthermore, we contrast the GW signals of AIC and rotating massive star iron core collapse and find that they can be distinguished, but only if the distance to the source is known and a detailed reconstruction of the GW time series from detector data is possible.  Some of our AIC models form massive quasi-Keplerian accretion disks after bounce. The disk mass is very sensitive to progenitor mass and angular momentum distribution. In rapidly differentially rotating models whose precollapse masses are significantly larger than the Chandrasekhar mass, the resulting disk mass can be as large as $ \\sim 0.8 M_\\odot $.  Slowly and/or uniformly rotating models that are limited to masses near the Chandrasekhar mass produce much smaller disks or no disk at all. Finally, we find that the postbounce cores of rapidly spinning white dwarfs can reach sufficiently rapid rotation to develop a gravito-rotational bar-mode instability. Moreover, many of our models exhibit sufficiently rapid and differential rotation to become subject to recently discovered low-$ E_\\mathrm{rot} / |W| $-type dynamical instabilities. ",
        "introduction": "\\label{sec:introduction}  Single stars with main-sequence masses $M \\lesssim 100\\, \\mmsun$ end their nuclear-burning lives as electron-degenerate objects or with central electron-degenerate cores. More specifically, the end state is a carbon-oxygen or oxygen-neon white dwarf (WD) in the case of low-mass stars (\\ie with $M \\lesssim 6-8\\, \\mmsun$), or a degenerate oxygen-neon or iron core embedded in an extended non-degenerate stellar envelope in the case of more massive stars (\\ie $6-8\\, \\mmsun \\lesssim M \\lesssim 100\\,\\mmsun$) (see, \\eg \\cite{herwig:05,woosley_02_a,poelarends:08} and references therein). Electron-degenerate spherically-symmetric objects become unstable to radial contraction once their mass exceeds the Chandrasekhar mass which, assuming zero temperature and no rotation, is given by $M_{\\mathrm{Ch}} = 1.4575\\, (Y_e/0.5)^2\\,\\mmsun\\,$, where $Y_e$ is the number of electrons per baryon, or ``electron fraction''~\\cite{chandrasekhar:38,shapteu:83}. The \\emph{effective} Chandrasekhar mass $M_{\\mathrm{Ch,eff}}$ of a WD or a stellar core increases somewhat with WD/core entropy (\\eg \\cite{woosley_02_a}) and can grow considerably by rotation, in which case it is limited only by the onset of nonaxisymmetric instability (\\eg \\cite{ostriker_68_b,yoon_04,shapteu:83}).  The iron core of a massive star is pushed over its Chandrasekhar limit by the ashes of silicon shell burning and undergoes collapse to a protoneutron star (PNS), accelerated by photodisintegration of heavy nuclei and electron capture \\cite{bethe:90}.  In a high-density sub-$M_\\mathrm{ch}$ oxygen-neon core of a less massive star, electron capture may decrease $M_\\mathrm{Ch,eff}$, also leading to collapse \\cite{nomoto_84_a,gutierrez_05_a}.  In both cases, if an explosion results, the observational display is associated with a Type II/Ibc supernova (SN).   On the other hand, a carbon-oxygen WD can be pushed over its stability limit through merger with or accretion from another WD (double-degenerate scenario) or by accretion from a non-degenerate companion star (single-degenerate scenario). Here, the WD generally experiences carbon ignition and thermonuclear runaway, leading to a Type Ia SN and leaving no compact remnant~\\cite{livio_99_a}.  However, at least theoretically, it is possible that massive oxygen-neon WDs\\footnote{Previously, such WDs were expected to have a significant central $^{24}{\\mathrm{Mg}}$ mass fraction, hence were referred to as oxygen-neon-magnesium WDs. Recent work based on up-do-date input physics and modern stellar evolution codes suggests that the mass fraction of $^{24}{\\mathrm{Mg}}$ is much smaller than previously thought (\\eg~\\cite{siess:06,iben:97,ritossa:96}).} formed by accretion or merger, and, depending on initial mass, temperature, and accretion rate, also carbon-oxygen WDs, may grow to reach their $M_{\\mathrm{Ch,eff}}$ or reach central densities sufficiently high ($\\gtrsim 10^{9.7} - 10^{10}\\,{\\mathrm{g\\,cm^{-3}}}$) for rapid electron capture to take place, triggering collapse to a PNS rather than thermonuclear explosion~\\cite{canal:76,saio:85,nomoto:86,mochkovitch_89_a,nomoto_91,saio:98, uenishi_03_a,saio_04,yoon_04,yoon_05,yoon_07_a,kalogera_00_a, gutierrez_05_a}. This may result in a peculiar, in most cases probably sub-energetic, low-nickel-yield and short-lived transient~\\cite{nomoto:86,woosley_92_a,fryer_99_a,dessart_06_a,dessart_07_a,metzger:09b}. This alternative to the Type Ia SN scenario is called \\textit{``accretion-induced collapse''} (AIC) and will be the focus of this paper.  The details of the progenitor WD structure and formation and the fraction of all WDs that evolve to AICs are presently uncertain. Binary population synthesis models \\cite{yungelson_98_a,belczynski_05_a, kalogera_00_a} and constraints on $r$-process nucleosynthetic yields from previous AIC simulations \\cite{fryer_99_a,qian:07} predict AIC to occur in the Milky Way at a frequency of $\\sim 10^{-5}$ to $\\sim10^{-8}\\,{\\mathrm{yr}}^{-1}$ which is $\\sim 20-50$ times less frequent than the expected rate of standard Type Ia SNe~(\\eg~\\cite{vdb:91,madau:98,scannapieco:05,mannucci:05}). In part as a consequence of their rarity, but probably also due to their short duration and potentially weak electromagnetic display, AIC events have not been directly observed (but see \\cite{perets:09}, who discovered a peculiar type-Ib SN in NGC 1032 that can be interpreted as resulting from an AIC).  The chances of seeing a rare galactic AIC are dramatically boosted by the possibility of guiding electromagnetic observations by the detection of neutrinos and gravitational waves (GWs) emitted during the AIC process and a subsequent SN explosion.  GWs, similar to neutrinos, are extremely difficult to observe, but can carry ``live'' dynamical information from deep inside electromagnetically-opaque regions. The inherent multi-D nature of GWs (they are lowest-order quadrupole waves) makes them ideal messengers for probing multi-D dynamics such as rotation, turbulence, or NS pulsations~\\cite{thorne:87,andersson:03,ott:09rev}.  The detection prospects for a GW burst from an AIC are significantly enhanced if theoretical knowledge of the expected GW signature of such an event is provided by computational modelling. In reverse, once a detection is made, detailed model predictions will make it possible to extract physical information on the AIC dynamics and the properties of the progenitor WD and, hence, will allow ``parameter-estimation'' of the source.  Early spherically-symmetric (1D) simulations of AIC~\\cite{mayle_88_a,baron_87_b,woosley_92_a} and more recent axisymmetric (2D) ones~\\cite{fryer_99_a,dessart_06_a, dessart_07_a} have demonstrated that the dynamics of AIC is quite similar to standard massive star core collapse: During collapse, the WD separates into a subsonically and homologously collapsing ($v \\propto r$) inner core and a supersonically collapsing outer core. Collapse is halted by the stiffening of the equation of state (EOS) at densities near nuclear matter density and the inner core rebounds into the still infalling outer core. An unshocked low-entropy PNS of inner-core material is formed. At its edge, a bounce shock is launched and initially propagates rapidly outward in mass and radius, but loses energy to the dissociation of heavy nuclei as well as to neutrinos that stream out from the optically thin postshock region. The shock stalls and, in the AIC case (but also in the case of the oxygen-neon core collapse in super-AGB stars \\cite{kitaura_06_a,burrows:07c}), is successfully revived by the deposition of energy by neutrinos in the postshock region (\\ie the ``delayed-neutrino mechanism''~\\cite{bethewilson:85,bethe:90}) or by a combination of neutrino energy deposition and magnetorotational effects in very rapidly rotating WDs~\\cite{dessart_07_a}.  But even without shock revival, explosion would occur when the WD surface layer is eventually accreted through the shock.  Following the onset of explosion, a strong long-lasting neutrino-driven wind blows off the PNS surface, adding to the total explosion energy and establishing favorable conditions for $r$-process nucleosynthesis~\\cite{woosley_92_a,fryer_99_a, dessart_06_a,dessart_07_a,arcones:07}. If the progenitor WD was rotating rapidly (and had a rotationally-enhanced $M_\\mathrm{Ch,eff}$), a quasi-Keplerian accretion disk of outer-core material may be left after the explosion~\\cite{dessart_06_a}. Metzger~et~al.~\\cite{metzger:09,metzger:09b} recently proposed that this may lead to nickel-rich outflows that could significantly enhance the AIC observational display.    Rotating iron core collapse and bounce is the most extensively studied and best understood GW emission process in the massive star collapse context (see, \\eg~\\cite{dimmelmeier_08_a} and the historical overview in~\\cite{ott:09rev}). However, most massive stars (perhaps up to $\\sim 99\\%$ in the local universe) are likely to be rather slow rotators that develop little asphericity during collapse and in the early postbounce phase~\\cite{heger:05,woosley:06b,ott:06spin} and produce PNSs that cool and contract to neutron stars with periods above $\\sim 10\\,\\mathrm{ms}$ and parameter $\\beta = E_\\mathrm{rot}/|W|\\lesssim 0.1\\%$ \\cite{ott:06spin}, where $E_\\mathrm{rot}$ is the rotational kinetic energy and $|W|$ is the gravitational binding energy.  This does not only reduce the overall relevance of this emission process, but also diminishes the chances for postbounce gravito-rotational nonaxisymmetric deformation of the PNS which could boost the overall GW emission~\\cite{ott:09rev}. Axisymmetric rapidly rotating stars become unstable to nonaxisymmetric deformations if a nonaxisymmetric configuration with a lower total energy exists at a given $ \\beta $ (see~\\cite{stergioulas:03} for a review).  The classical high-$\\beta$ instability develops in Newtonian stars on a dynamical timescale at $\\beta \\gtrsim \\beta_\\mathrm{dyn} \\simeq 27\\%$ (the general-relativistic value is $\\beta \\gtrsim 25\\%$ \\cite{baiotti_07_a,Manca07}). A ``secular'' instability, driven by fluid viscosity or GW backreaction, can develop already at $ \\beta \\gtrsim \\beta_{\\mathrm{sec}} \\simeq 14\\% $ \\cite{stergioulas:03}. Slower, but strongly differentially rotating stars may also be subject to a nonaxisymmetric dynamical instability at $ \\beta $ as small as $ \\sim 1\\% $. This instability at low $\\beta$ was observed in a number of recent 3D simulations (\\eg \\cite{centrella_01_a, shibata:04a, ott_05_a, ou_06_a, ott:06phd, cerda_07_b, ott_07_a, ott_07_b,scheidegger:08}), and may be related to corotation instabilities in disks, but its nature and the precise conditions for its onset are presently not understood \\cite{watts:05,saijo:06}.  Stellar evolution theory and pulsar birth spin estimates suggest that most massive stars are rotating rather slowly (\\eg \\cite{heger:05, ott:06spin}, but also \\cite{cantiello:07,woosley:06b} for exceptions). Hence, rotating collapse and bounce and nonaxisymmetric rotational instabilities are unlikely to be the dominant GW emission mechanisms in most massive star collapse events \\cite{ott:09rev}. The situation may be radically different in AIC: Independent of the details of their formation scenario, AIC progenitors are expected to accrete significant amounts of mass and angular momentum in their pre-AIC evolutions \\cite{yoon_04,saio_04,yoon_05,uenishi_03_a, piersanti_03_a}. They may reach values of $\\beta$ of up to $\\sim 10\\%$ \\emph{prior} to collapse, according to the recent work of Yoon \\& Langer~\\cite{yoon_04,yoon_05}, who studied the precollapse stellar structure and rotational configuration of WDs with sequences of 2D rotational equilibria.  Depending on the distribution of angular momentum in the WD, rotational effects may significantly affect the collapse and bounce dynamics and lead to a large time-varying quadrupole moment of the inner core, resulting in a strong burst of GWs emitted at core bounce. In addition, the postbounce PNS may be subject to the high-$\\beta$ rotational instability (see \\cite{liu_01_a,liu_02_a} for an investigation via equilibrium sequences of PNSs formed in AIC) or to the recently discovered low-$\\beta$ instability.  Most previous (radiation-)hydrodynamic studies of AIC have either been limited to 1D~\\cite{mayle_88_a,baron_87_b,woosley_92_a} or were 2D, but did not use consistent 2D progenitor models in rotational equilibrium~\\cite{fryer_99_a}.  Fryer, Hughes and Holz~\\cite{fryer:02} presented the first estimates for the GW signal emitted by AIC based on one model of~\\cite{fryer_99_a}. Drawing from the Yoon \\& Langer AIC progenitors~\\cite{yoon_04,yoon_05}, Dessart~et~al.~\\cite{dessart_06_a,dessart_07_a} have recently performed 2D Newtonian AIC simulations with the multi-group flux-limited diffusion (MGFLD) neutrino radiation-(M)HD code VULCAN/2D \\cite{livne:93,livne:07,burrows:07a}. They chose two representative WD configurations for slow and rapid rotation with central densities of $ 5 \\times 10^{10} \\ {\\mathrm{g\\,cm}}^{-3} $ and total masses of $ 1.46 M_\\odot $ and $ 1.92 M_\\odot $. Both models were set up with the differential rotation law of \\cite{yoon_04,yoon_05}. The $1.46 M_\\odot$ model had zero rotation in the inner core and rapid outer core rotation while the $1.92 M_\\odot$ was rapidly rotating throughout (ratio $\\Omega_\\mathrm{max,initial} / \\Omega_\\mathrm{center,initial} \\sim 1.5$). Dessart~et~al.~\\cite{dessart_06_a,dessart_07_a} found that rapid electron capture in the central regions of both models led to collapse to a PNS within only a few tens of milliseconds and reported successful neutrino-driven~\\cite{dessart_06_a} and magnetorotational explosions~\\cite{dessart_07_a} with final values of $\\beta$ (\\ie a few hundred milliseconds after core bounce) of $\\sim 6\\%$ and $\\sim 26\\%$, for the $1.46 \\mmsun$ and $1.92 \\mmsun$ models, respectively\\footnote{These numbers are for the non-MHD simulations of~\\cite{dessart_06_a}. In the MHD models of ~\\cite{dessart_07_a}, an $\\Omega$-dynamo builds up toroidal magnetic field, reducing the overall rotational energy and $\\beta$.}. The analysis in \\cite{dessart_06_a,ott:06phd,ott:09rev} of the GW signal of the Dessart~et~al.\\ models showed that the morphology of the AIC rotating collapse and bounce gravitational waveform is reminiscent of the so-called Type~III signal first discussed by Zwerger \\& M\\\"uller~\\cite{zwerger_97_a} and associated with small inner core masses and a large pressure reduction at the onset of collapse in the latter's polytropic models.  In this paper, we follow a different approach than Dessart~et~al.~\\cite{dessart_06_a,dessart_07_a}. We omit their detailed and computationally-expensive treatment of neutrino radiation transport in favor of a simple, yet effective deleptonization scheme for the collapse phase \\cite{liebendoerfer_05_a}. This simplification, while limiting the accuracy of our models at postbounce times $\\gtrsim 5-10\\,\\mathrm{ms}$, (i) enables us to study a very large set of precollapse WD configurations and their resulting AIC dynamics and GW signals and, importantly, (ii) allows us to perform these AIC simulations in \\emph{general relativity}, which is a crucial ingredient for the accurate modeling of dynamics in regions of strong gravity inside and near the PNS. Furthermore, as demonstrated by \\cite{dimmelmeier_02_b,dimmelmeier_07_a,dimmelmeier_08_a}, general relativity is required for qualitatively and quantitatively correct predictions of the GW signal of rotating core collapse.  We focus on the collapse and immediate postbounce phase of AIC and perform an extensive set of \\nummodel\\ 2D general-relativistic hydrodynamics simulations.  We analyze systematically the AIC dynamics and the properties of the resulting GW signal. We explore the dependence of nonrotating and rotating AIC on the precollapse WD rotational setup, central density, core temperature, and core deleptonization, and study the resulting PNS's susceptibility to rotational nonaxisymmetric deformation.  Furthermore, motivated by the recent work of Metzger~et~al.\\ \\cite{metzger:09,metzger:09b}, who discussed the possible enhancement of the AIC observational signature by outflows from PNS accretion disks, we study the dependence of disk mass and morphology on WD progenitor characteristics and rotational setup.  We employ the general-relativistic hydrodynamics code \\coconut ~\\cite{dimmelmeier_02_b,dimmelmeier_05} and neglect MHD effects since they were shown to be small in the considered phases unless the precollapse magnetic field strength is extremely large ($B \\gtrsim 10^{12}\\, {\\mathrm G}$, \\eg \\cite{obergaulinger_06_a,dessart_07_a,burrows_07_b}).  We employ a finite-temperature microphysical nuclear EOS in combination with the aforementioned deleptonization treatment of \\cite{liebendoerfer_05_a}. The precollapse 2D rotational-equilibrium WDs are generated according to the prescription of Yoon \\& Langer~\\cite{yoon_04,yoon_05}.  The plan of the paper is as follows. In Sec.~\\ref{sec:methods}, we introduce the numerical methods employed and discuss the generation of our 2D rotational-equilibrium precollapse WD models as well as the parameter space of WD structure and rotational configuration investigated.  In Sec.~\\ref{sec:colldyn}, we discuss the overall AIC dynamics and the properties of the quasi-Keplerian accretion disks seen in many models. Sec.~\\ref{sec:GW} is devoted to a detailed analysis of the GW signal from rotating AIC. There, we also assess the detectability by current and future GW observatories and carry out a comparison of the GW signals of AIC and massive star iron core collapse.  In Sec.~\\ref{sec:rotinst}, we study the postbounce rotational configurations of the PNSs in our models and assess the possibility for nonaxisymmetric rotational instabilities. In Sec.~\\ref{sec:summary}, we present a critical summary and outlook. ",
        "conclusions": "\\label{sec:summary}   In this paper we have presented the first general-relativistic simulations of the axisymmetric AIC of massive white dwarfs to protoneutron stars. Using the general-relativistic hydrodynamics code \\coconut, we performed \\nummodel\\ baseline model calculations, each starting from a 2D equilibrium configuration, using a finite-temperature microphysical EOS, and a simple, yet effective parametrization scheme of the electron fraction $Y_e$ that provides an approximate description of deleptonization valid in the collapse, bounce, and very early postbounce phases.  The precollapse structure and rotational configuration of WDs that experience AIC is essentially unconstrained. This prompted us to carry out this work. With our large set of model calculations, we have investigated the effects on the AIC evolution of variations in precollapse central density, temperature, central angular velocity, differential rotation, and deleptonization in collapse.  The inclusion of general relativity enabled us to correctly describe the AIC dynamics and our extended model set allowed us for the first time to study systematically GW emission in the AIC context.  We find that the overall dynamics in the collapse phase of AIC events is similar to what has long been established for rotating iron core collapse. A universal division in homologously collapsing inner core and supersonically infalling outer core obtains and the self-similarity of the collapse nearly completely washes out any precollapse differences in stellar structure in the limit of slow rotation. Due to the high degeneracy of the electrons in the cores of AIC progenitor WDs, electron capture is predicted to be strong already in early phases of collapse \\cite{dessart_06_a}, leading to a low trapped lepton fraction and consequently small inner core masses $M_\\mathrm{ic,b}$ at bounce of around $0.3\\,M_\\odot$ which decrease somewhat with increasing precollapse WD temperature due to the temperature dependent abundance of free protons.  Test calculations motivated by potential systematic biases of the AIC $\\overline{Y_e}(\\rho)$ trajectories obtained from \\cite{dessart_06_a} (see Secs.~\\ref{sec:deleptonization} and \\ref{sec:tempye}) with inner-core values of $Y_e$ increased by $\\sim 10\\%$ and $\\sim 20\\%$ yielded values $M_\\mathrm{ic,b}$ larger by $\\sim 11\\%$ and $\\sim 25\\%$.  Our simulations show that rotation can have a profound influence on the AIC dynamics, but will \\emph{always} stay subdominant in the collapse of uniformly rotating WDs whose initial angular velocity is constrained by the Keplerian limit of surface rotation. In rapidly differentially rotating WDs, on the other hand, centrifugal support can dominated the plunge phase of AIC and lead to core bounce at subnuclear densities. We find that the parameter $\\beta_\\mathrm{ic,b} = (E_\\mathrm{rot}/|W|)_\\mathrm{ic,b}$ of the inner core at bounce provides a unique mapping between inner core rotation and late-time collapse and bounce dynamics, but the mapping between precollapse configurations and $\\beta_\\mathrm{ic,b}$ is highly degenerate, i.e., multiple, in many cases very different precollapse configurations of varying initial compactness and total angular momentum, can yield practically identical $\\beta_\\mathrm{ic,b}$ and corresponding collapse/bounce dynamics.  Recent phenomenological work presented in \\cite{metzger:09,metzger:09b} on the potential EM display of an AIC event has argued for both uniform WD rotation \\cite{metzger:09,piro_08_a} and massive quasi-Keplerian accretion disks left behind at low latitudes after AIC shock passage. The analysis of our extensive model set, on the other hand, shows that uniformly rotating WDs produce no disks at all or, in extreme cases that are near mass shedding at the precollapse stage, only very small disks ($M_\\mathrm{disk} \\lesssim 0.03 M_\\odot$). Only rapidly differentially rotating WDs yield the large disk masses needed to produce the enhanced EM signature proposed in \\cite{metzger:09,metzger:09b}.  An important focus of this work has been on the GW signature of AIC. GWs, due to their inherently multi-D nature, are ideal messengers for the rotational dynamics of AIC.  We find that all AIC models following our standard $\\overline{Y_e}(\\rho)$ parametrizations yield GW signals of a generic morphology which has been classified previously as \\emph{type~III} \\cite{zwerger_97_a,dimmelmeier_02_b,ott:06phd}.  This signal type is due primarily to the small inner core masses at bounce obtained in these models.  We distinguish between three subtypes of AIC GW signals. Type IIIa occurs for $\\beta_\\mathrm{ic,b} \\lesssim 0.7\\%$ (slow rotation), is due in part to early postbounce prompt convection and results in peak GW amplitudes $|h_\\mathrm{max}| \\lesssim 5 \\times 10^{-22}$ (at $10\\,\\mathrm{kpc}$) and emitted energies $E_\\mathrm{GW} \\lesssim \\mathrm{few}\\,\\times\\,10^{-9}\\,M_\\odot \\, c^2$. Most of our AIC models produce type~IIIb GW signals that occur for $0.7\\lesssim \\beta_\\mathrm{ic,b} \\lesssim 18\\%$ (moderate/moderately rapid rotation) and yield $6\\times10^{-22} \\lesssim |h_\\mathrm{max}|\\,(\\mathrm{at}\\,10\\,\\mathrm{kpc}) \\lesssim 8\\times10^{-21}$ and emitted energies of $9\\times10^{-10} M_\\odot\\,c^2 \\lesssim E_\\mathrm{GW} \\lesssim 2\\times10^{-8}\\,M_\\odot\\,c^2$. Rotation remains subdominant in type~IIIa and type~IIIb models and we find that there is a monotonic and near-linear relationship between maximum GW amplitude and the rotation of the inner core which is best described by the power law $|h_\\mathrm{max}| \\propto 10^{-21} \\beta_\\mathrm{ic,b}^{0.74}$. Furthermore, we find that the frequencies $f_\\mathrm{max}$ at which the GW spectral energy densities of type IIIa and IIIb models peak are in a rather narrow range from $\\sim 720\\,\\mathrm{Hz}$ to $\\sim 840\\,\\mathrm{Hz}$ and exhibit a monotonic growth from the lower to the upper end of this range with increasing rotation. This finding suggests that the GW emission in these models is driven by the fundamental quadrupole (${}^2\\!f$) mode of the inner core.  In the dynamics of AIC models that reach $\\beta_\\mathrm{ic,b} \\gtrsim 18\\%$, centrifugal effects become dominant and lead to core bounce at subnuclear densities. Such models must be differentially rotating at the onset of collapse and produce type~IIIc GW signals with maximum amplitudes of $4.0\\times10^{-22} \\lesssim |h_\\mathrm{max}|\\,\\mathrm{(at\\, 10\\,\\mathrm{kpc})} \\lesssim 5.5\\times10^{-21}$, emitted energies of $ 10^{-10} M_\\odot\\,c^2 \\lesssim E_\\mathrm{GW} \\lesssim 10^{-8} M_\\odot\\, c^2$, and peak frequencies of $62\\, \\mathrm{Hz} \\lesssim f_\\mathrm{max} \\lesssim 800 \\, \\mathrm{Hz}$. In contrast too type~IIIa and IIIb models, in type~IIIc models, $|h_\\mathrm{max}|$, $E_\\mathrm{GW}$, and $f_\\mathrm{max}$ decrease monotonically with increasing $\\beta_\\mathrm{ic,b}$.   Combining the information from signal morphology, $|h_\\mathrm{max}|$, $E_\\mathrm{GW}$ and $f_\\mathrm{max}$, we conclude that already first-generation interferometer GW detectors should be able to infer the rotation of the inner core at bounce (as measured by $\\beta_\\mathrm{ic,b}$) from a Galactic AIC event. Due to the degenerate dependence of $\\beta_\\mathrm{ic,b}$ on initial model parameters, this can put only loose constraints on the structure and rotational configuration of the progenitor WD.  However, the observation of an AIC with $\\beta_\\mathrm{ic,b} \\gtrsim 18\\%$ would rule out uniform progenitor rotation.     Studying the configurations of the protoneutron stars formed in our AIC models, we find that none of them are likely to experience the high-$\\beta$ nonaxisymmetric bar-mode instability at very early postbounce times. We estimate, however, that all models that reach $\\beta_\\mathrm{ic,b} \\gtrsim 17\\%$ will contract and reach the instability threshold within $\\sim 50\\,\\mathrm{ms}$ after bounce. Less rapidly spinning models will require more time or will go unstable to the low-$\\beta$ instability. The latter requires strong differential rotation which is ubiquitous in the outer PNS and in the postshock region of our AIC models.  AIC progenitors, due to their evolution through accretion or formation through merger, are predestined to be rapidly rotating and form PNSs that are likely to become subject to nonaxisymmetric instabilities. This is in contrast to the precollapse iron cores of ordinary massive stars that are expected to be mostly slowly-spinning objects~\\cite{heger:05,ott:06spin}. We conclude that the appearance of nonaxisymmetric dynamics driven by either the low-$\\beta$ or high-$\\beta$ instability and the resulting great enhancement of the GW signature may be a generic aspect of AIC and must be investigated in 3D models.  The comparison of the GW signals of our axisymmetric AIC models with the gravitational waveforms of the iron core collapse models of Dimmelmeier~et~al.~\\cite{dimmelmeier_08_a} reveals that the overall characteristics of the signals are rather similar. It appears unlikely that AIC and iron core collapse could be distinguished on the basis of the axisymmetric parts of their GW signals alone, unless detailed knowledge of the signal time series as well as of source orientation and distance is available to break observational degeneracies.  The results of our AIC simulations presented in this paper and the conclusions that we have drawn on their basis demonstrate the complex and in many cases degenerate dependence of AIC outcomes and observational signatures on initial conditions. The observation of GWs from an AIC event can provide important information on the rotational dynamics of AIC.  However, to lift degeneracies in model parameters and gain full insight, GW observations must be complemented by observations of neutrinos and electromagnetic waves.  These multi-messenger observations require underpinning by comprehensive and robust computational models that have no symmetry constraints and include all the necessary physics to predict neutrino, electromagnetic, and GW signatures.   As a point of caution, we note that the generic type~III GW signal morphology observed in our AIC models is due to the small inner-core values of $Y_e$ and consequently small inner core masses predicted by the $\\overline{Y_e}(\\rho)$ parametrization obtained from the approximate Newtonian radiation-hydrodynamic simulations of \\cite{dessart_06_a}.  Tests with artificially reduced deleptonization show that the signal shape becomes a mixture of type~III found in our study and type I observed in rotating iron core collapse \\cite{dimmelmeier_08_a} if the $Y_e$ in the inner core is larger by $\\sim 20\\%$.  In a follow-up study, we will employ $\\overline{Y_e}(\\rho)$ data from improved general-relativistic radiation-hydrodynamics simulations \\cite{mueller:09phd} to better constrain the present uncertainties of the AIC inner-core electron fraction.  Although performed using general-relativistic hydrodynamics, the calculations discussed here are limited to conformally-flat spacetimes and axisymmetry. We ignored postbounce deleptonization, neutrino cooling, and neutrino heating.  We also neglected nuclear burning, employed only a single finite-temperature nuclear EOS, and were forced to impose ad-hoc initial temperature and electron fraction distributions onto our precollapse WD models in rotational equilibrium.  Future studies must overcome the remaining limitations to build accurate models of AIC.  Importantly, extensive future 3D radiation-hydrodynamic simulations are needed to address the range of possible, in many cases probably nonaxisymmetric, postbounce evolutions of AIC and to make detailed predictions of their signatures in GWs, neutrinos, and in the electromagnetic spectrum."
    },
    "0910/0910.2190_arXiv.txt": {
        "abstract": "We perform population synthesis studies of different types of neutron stars (thermally emitting isolated neutron stars, normal radio pulsars, magnetars) taking into account the magnetic field decay and using results from the most recent advances in neutron star cooling theory. For the first time, we confront our results with observations using {\\it simultaneously} the Log N -- Log S distribution for nearby isolated neutron stars, the Log N -- Log L distribution for magnetars, and the distribution of radio pulsars in the $P$ -- $\\dot P$ diagram. For this purpose, we fix a baseline neutron star model (all microphysics input), and other relevant parameters to standard values (velocity distribution, mass spectrum, birth rates ...), allowing to vary the initial magnetic field strength. We find that our theoretical model is consistent with all sets of data if the initial magnetic field distribution function follows a log-normal law with $<\\log (B_0/[G])>\\sim 13.25$ and $\\sigma_{\\log B_0}\\sim 0.6$. The typical scenario includes about 10\\% of neutron stars born as magnetars, significant magnetic field decay during the first million years of a NS life (only about a factor of 2 for low field neutron stars but more than an order of magnitude for magnetars), and a mass distribution function dominated by low mass objects. This model explains satisfactorily all known populations. Evolutionary links between different subclasses may exist, although robust conclusions are not yet possible. ",
        "introduction": "At the present moment our knowledge about neutron star (NS) evolution is an intriguing puzzle. We know many observational manifestations of young isolated NSs: radio pulsars (PSRs); central compact objects in supernova remnants (CCOs in SNRs);  rotating radio transients (RRATs); radio-quiet thermally emitting isolated NSs, also known as X-ray dim isolated NSs (XDINS) or the Magnificent Seven (M7), and the observational manifestations of magnetars --soft gamma-ray repeaters (SGRs) and anomalous X-ray pulsars (AXPs). Reasons for this apparent diversity as well as possible links between different classes are not entirely clear (see e.g. \\cite{psr2005,WT2006,Kaspi2007,Haberl2007,Zane2007,rrat2008} for recent reviews about the different subclasses).  In the past few years we have learned that some NSs can show different types of activity, transiting from class to class. For example, PSR J1846-0258, known for some time as a normal pulsar, demonstrated outbursts typical for AXPs or/and SGRs \\citep{kumar2008, gavriil2008}. In addition, the total energy release by the object became larger than the rotational energy losses, which according to the original classification discussed in \\cite{TD1996} should place it in the AXP list. Thus, for the first time, we have an example of transformation of a PSR into a magnetar. Several of the SGRs do not show any bursting activity for many years, and if we had not enough information about their violent past, we would have classified them as AXPs. Some of RRATs were shown to emit normal radio pulsar emission \\citep{deneva2009,rrats2009}. The transient AXP XTE J1810-197 and AXP 1E 1547.0-5408 demonstrated radio pulses \\citep{camilo2006, camilo2007}. One of the RRATs J1819-1458 shows thermal properties much similar to the M7 \\citep{reynolds2006}. Hence, divisions between some subpopulations of young isolated NSs (or at least some of their representatives) can be illusive. On the other hand, young, low-field CCOs \\citep{halpern2007}, normal PSRs ($B\\sim 10^{12}$~G, i.e. Crab and Vela-like), and SGRs clearly represent NSs born with different properties.  Our brief observational record of NSs ($\\sim 40$ years at most), low statistics in many cases, and selection effects do not allow to draw a coherent picture of NS ``sociology'' just from observations. From the theoretical point of view, our understanding of the SN explosion mechanism is not precise enough to provide a solid model of initial parameters of NSs  and evolutionary models (thermal, magnetic field, and spin evolution) are related to extremely complicated physical problems (superfluidity and superconductivity in dense matter, electrodynamics in superstrong magnetic fields, etc.) which usually leads to inconclusive results.  In our opinion, it is necessary not only to compare observations with models to verify individual objects, but also to confront theoretical calculations with observational data via population synthesis techniques taking into account as many classes of NSs as possible. Joint constraints by means of simultaneous comparison of theoretical models with different subpopulations should be derived to form a {\\it population mosaic}. Numerous degrees of freedom in modern evolutionary models must be compensated by different observational tests. Several important studies in this area appeared in recent years, see for example \\cite{fgkaspi2006, keane2008}, and references therein. In this paper we continue to follow this lane but with an important difference: we attempt to constrain our model and check its consistency using at the same time different populations. This is, to our knowledge, the first time that a multilateral approach is employed. Previous works that focused in a specific NS observational class, ignored that the parameters obtained may be in contradiction with the properties of a different class of objects. For examples, models without magnetic field decay are clearly in contradiction with the existence of magnetars. We try to give a step forward in the direction towards a NS unified model by using at the same time different populations.  Among the different physical ingredients needed to properly model the thermal evolution of NSs, we emphasize that heat transport in the NS crust plays a crucial role during the first million years of its thermal evolution \\citep{Yakovlev2004}, which is typically the period during which NSs are detectable with current  X-ray instruments. Multi-wavelength observations in the soft X-ray, UV, and optical bands of the thermal emission from a NS's surface now provide a real opportunity to probe the internal physics of NSs (see \\citealt{Page2006} for a general overview).  Remarkably, all isolated nearby compact X-ray sources that have been detected also in the optical band (RXJ 185635-3754, RX J0720.4-3125, RX J1308.6+2127, and RX J1605.3+3249) have a significant optical excess relative to the extrapolated X-ray blackbody emission \\citep{Haberl2007}. An  inhomogeneous surface temperature distribution can accommodate this optical excess and can arise naturally if heat conduction in the NS crust is anisotropic due to the presence of a large magnetic field \\citep{Geppert2004,Geppert2006,Azorin2006a}. Alternative models are based on anisotropic radiation from magnetized atmospheres, as in \\citet{Ho2007}, or condensed surfaces as in \\citet{Azorin2005}. None of the explanations (only temperature anisotropy vs. physical origin) are entirely satisfactory and probably a combination of both effects is needed. Heat conduction can also influence other observable aspects of accreting NSs in low mass X-ray binaries, including their quiescent luminosity and the superburst recurrence time-scales \\citep{Brown1998}. For this reason, one of the goals of this paper is to revisit former population synthesis studies using new NS evolution models that include anisotropic heat transport in NSs crusts and magnetic field evolution \\citep{Aguilera2008b,Aguilera2008a,PMG2009}.  In this paper we want to study if the {\\it patchy} view of young NSs subpopulations can be explained by a unique set of smooth distributions of the most important parameters, among which the the magnetic field distribution plays the main role.\\footnote{Note for the arXiv version: After the final (second revision) version of this paper has been submitted, an important e-print by Kaplan and van Kerkwijk appeared (arXiv: 0909.5218). In this article the authors give new observational arguments in favor of the large role of the field decay in the evolution of the Magnificent Seven. Many of our conclusions coincide with those given by Kaplan and van Kerkwijk.} Our main idea here is to  use several different tests to confront theoretical predictions and observations. For this purpose, the paper is organized as follows. In the next section we briefly describe the model of magneto-thermal evolution of NSs that we use. With this model, one has the cooling behavior and magnetic field (and therefore period) evolution of NS. In Sec.~3, we describe the population synthesis technique used for Log N~--~Log S calculations of close-by cooling NSs which thermal emission has been detected and which temperature has been estimated. With this, we can partially constrain the initial magnetic field distribution. Then, in Sec.~4, we discuss the Log N~--~Log L distribution of the population of magnetars in the Galaxy. We show that the model constrained in the previous section is also consistent with the flux distributions of the extrapolated magnetar population. To finish the presentation of our results, in Sec.~5 we use the population of radio pulsars in the $P-\\dot P$ diagram to put additional constraints on the properties of NSs. Explicitly, the current period and magnetic field distributions of rotation-powered pulsars constrain the NS initial magnetic field distribution and break the degeneracy in the parameter space obtained in previous section. Sec.~6 is devoted to the final remarks  and to discuss uncertainties of the models we use and  future prospects. ",
        "conclusions": "In this paper we presented a multi-component population synthesis study. The final goal in this approach would be to make a complete population synthesis with a unique NS physical model that consistently explains all known types of young NSs just varying parameters such as the NS mass, age, or the strength and geometry of the magnetic field. Still, even after more sophisticated theoretical calculations are available, comparison with the data will proceed by steps, confronting each piece of simulated data with observational data of some population, in a similar way to this paper. One reason for engaging multi-population studies can be illustrated as follows. From Fig. 3 it is visible that the Log N -- Log S distribution can be explained by a single field model (with moderate additional heating, in our framework). However, the Log N -- Log L distribution for magnetars, as well as the $P$-$\\dot P$ plot for radio pulsars, of course, cannot be explained by this model for {\\it different} reasons. In the first case case, because no bright magnetars will be observed; in the latter, because it will be impossible to reproduce the main part of the pulsar population.  To summarize, we believe that the approach we use is well motivated by a necessity to have a natural model without {\\it ad hoc} assumptions about different subpopulations. Note that here we tested a very  ``smooth'' model, in a sense that we do not put by hand initially distinct populations (magnetars, M7, PSRs, etc.). In each part of our study (Log N -- Log S for the M7,  Log N -- Log L for magnetars, $P$--$\\dot P$ for PSRs) we model ``just NSs'' using the same initial magnetic field distribution (which in our model is the main parameter), without specifying unique particular features for a given subpopulation, as it  is usually done (see e.g.  \\citealt{ptp2006,gh2007,keane2008}).  For example, XDINSs are not a separately defined class with initially distinct properties, but they just appear as a population coming out from a smooth unique initial distribution, and with the following features: \\begin{itemize} \\item The magnetic field in these sources has significantly decayed from larger initial values, so no magnetar-like activity is present; \\item Due to their large initial fields spin periods are long, so no radio pulsar activity is observed (may be due to narrow beams); \\item They are still young enough  for the magnetic field to be still decaying ($<10^6$ yrs, according to \\cite{PMG2009}) and heating the crust, so that they can be observed as relatively bright thermal sources. \\end{itemize}  In our model a typical M7-like source has an initial field of $B \\sim 10^{14}$~G. The period is $P\\sim 7$~sec at a true age of $\\sim 5 \\, 10^5$~yrs (spin-down age $\\sim (8-9) \\, 10^5$~yrs). At this time such an object has $B\\sim 4 \\, 10^{13}$~G and $T\\sim (6-7) \\, 10^5$~K. The additional energy being input by field decay in these sources can be estimated in each case by multiplying the average magnetic energy density ($\\sim  10^{26}$ erg~cm$^{-3}$), by the crust volume ($ \\sim 10^{18}$ cm$^3$) and dividing by the typical Ohmic decay timescale under such conditions (crust density and temperature), which is  $\\sim 10^{13}$ s. This gives roughly $10^{30-31}$ erg~s$^{-1}$ available from magnetic field decay. A fraction of this energy can be radiated in the form of neutrinos, but the energy reservoir to keep the star warm is still very important.  We tested average ages and typical parameters (periods, magnetic fields, etc.) of NSs which contribute to the range of Log N -- Log S between  0.1 and 10 cts~s$^{-1}$, where all observed sources are located. On average modeled NSs with inital fields 10$^{14}$--3 10$^{14}$~G which contribute a lot to this range have ages 2 10$^5$ -- 5 10$^5$ yrs and fields (at the moment of observation) $\\sim 7\\, 10^{13}$~G. Periods are distributed between $\\sim$2-20 sec, in correspondence with observed properties of the M7. Of course, some fraction of lower field NSs are also found in this range, in correspondence with observations of cooling radio pulsars (Vela, Geminga etc.). Detailed study of the modeled population in the range 0.1-10 cts~s$^{-1}$ is in progress, and will be presented elsewhere.  In any population synthesis approach inevitably it is necessary to make simplifications and to neglect some details. Partly, simplifications made in this study are justified by low statistics of known sources, partly by uncertain properties of objects under study. One of the main problems in a population synthesis study is related to possible correlations between different parameters. Some correlations are irrelevant for our purposes (i.e., the correlation between direction of velocity and spin axis) but others can be crucial. For example, in our study we assume that masses, initial magnetic fields and spin periods are not correlated. However, this is very uncertain for all subclasses of NSs.  If these three parameters are correlated our results may change. For magnetars it was proposed that they can have massive progenitors \\citep{muno2006}, and normally more massive progenitors are expected to produce massive NSs \\citep{whw2002}. Thus, we can expect a correlation between field and mass for magnetars but not for other NS classes. Such correlations, valid only for a not well identified part of a population, are very hard to confirm or to rule out. In Fig. \\ref{fig1} one can see that an increase in the value of the magnetic field is much more influential (on the thermal evolution) than changes in mass. That is why we think that moderate mass-field correlation in the case of magnetars has very little influence on our results. Other effects like rotation and mass-loss would complicate the situation even more.  We also neglect all possible effects related to the fact that a significant fraction of NSs are born in binary systems. Potentially, this can be important for magnetars if their large magnetic fields are generated by amplification due to rapid rotation of the protoNS. The precise mechanism of magnetar formation is still unknown. In some models \\citep{pp2006,bp2009} magnetars are born only in binary systems where a progenitor core was spun-up due to accretion or tidal interaction, similar to the main scenario for gamma-ray burst progenitors.  In our simplified standard scenario (it is clearly visible in Fig.\\ref{ppdot}) the M7 sources do not seem to be descendants of extreme magnetars, but the evolutionary link with some of the AXPs is possible. However, due to numerical limitations and the very different timescales, our model does not include the possibility of enhanced magnetic field dissipation during the fast initial Hall stage, nor transient phenomena that change a quiet X-ray emitter into an active bursting source, and back. Understanding the short term violent behavior of young magnetars may help to reconcile even the highest field objects with the M7.  One possibility to test this hypothesis is to look at the velocity distribution of these types of sources. Indeed, velocity is a good invariant on the time scale of $\\sim$1 Myr. If AXPs are rapidly moving objects, as it was popular to assume some years ago, but M7 have relatively small velocities, then one would conclude that the two populations are not related.  Interestingly, there is a recent measure of the transverse velocity of the magnetar XTE J1810-197 \\citep{v1810}. The measured velocity is slightly below the average for normal young neutron stars, indicating that the mechanism of magnetar birth need not lead to high NS velocities. On the other hand, recently \\cite{Motch2009} reported new velocity measurements for the M7 sources, and one of them (RX J1308.6+2127) appears to be a fast object. In the near future we expect that the measurement of proper motions of some of these sources with new observations will contribute to our understanding of their evolutionary link, if any.  In our study we also assumed that the magnetic field structure is similar for all NSs and we only rescale the normalization value for a fixed geometry. This is obviously an arbitrary choice. There are several issues related with the magnetic field geometry that need  further investigation. If the magnetic field is not supported by currents located in the crust, but instead by superconducting currents in the core, the field will dissipate on much longer timescales, and the heating mechanism we assume is less important. The amount of energy stored in an internal toroidal field, compared to the corresponding energy of the measured dipolar component is also unknown, and it is an important parameter that can significantly alter the results. We have chosen toroidal fields which maximum amplitude is a factor of 2 the dipole estimate (however, the energy stored on the toroidal field is only about a 10\\% because it is confined to a small volume),  in agreement with recent studies on MHD equilibrium configurations \\citep{mhd1,mhd2}, but this remains an open question. Some models we tried with toroidal fields one order of magnitude larger produced much hotter objects, and overpredicted the observed number of isolated NSs and magnetars, unless we reduce significantly the birth rate.  In this paper we have not attempted to adjust our cooling curves for low initial magnetic fields in such a way that the local population of normal PSRs with detected thermal emission is perfectly explained. Future more extensive calculations exploring other parameters (superfluid gaps, crust physical properties, etc.) can help to fit better the population of thermally emitting local NSs, and a better knowledge of the mass spectrum can be important for low-field stars, too. For now, we can work with the known 7 XDINS (for which field decay is probably important) and 4-5 normal close-by PSRs (i.e. lower field sources without additional heating). In fact, the lowest curve in Fig.\\ref{lnls_aver} predicts only 1-2 objects with flux larger than 0.1 cts~s$^{-1}$, while we know four near-by PSRs (Vela, Geminga, B1055--52, and B0656+14). The low field cooling curves are more sensitive to details of the equation of state, superfluid gaps, and neutrino emissivities (and the NS mass) than high field models, and changes of this parameters (i.e. the neutron gap in the core) can shift up or down the curves.  Ideally, the low field population should be used in a separate analysis to constrain parameters of the interior physics but, given the problem of very low statistics we are facing, it seems meaningless to split our populations in more subgroups at this stage. We made some tests to see if it can be done but, with the low statistics we manage and with the present day uncertainties about details of NS cooling, we do not think that an extensive investigation of microphysical parameters can contribute much to our understanding of NS properties. We certainly need more data before making strong cases in favor of particular cooling models.  Despite many attempts (see \\citealt{a2006,c2005} and references therein) the number of M7-like sources is not increasing significantly. More M7-like NSs can be found using deep XMM-Newton and Chandra observations. A good example is the new source found by \\cite{pires2009}. Another possibility is to have a sample of $\\gamma$-ray sources selected  by Fermi/GLAST, but they should be not M7-like, but PSRs with radio beams not pointing towards us (similar to the ``second Geminga''). The most promising way to  increase the number of M7-like sources is related to the future eROSITA instrument aboard Russian satellite Spectrum-X-Gamma. New sources are expected to be dimmer, younger, hotter and further away than the seven ROSAT sources \\citep{posselt2008}, and eROSITA (with non-truncated sensitivity at low energies) will be a perfect tool for them. For magnetars there is some hope to have many more candidates due to the MAXI instrument aboard the International Space Station \\citep{maxi2009}. Then, with increased statistics, it will be useful to come back to study a separate Log N -- Log S distribution of close-by cooling PSRs.  To summarize, the methodology we employ in this article looks very promising to constrain NS properties as more data coming from future missions arrive and more PSRs are found by radio surveys. A future increase of the observational data will allow to perform much more detailed population synthesis and to establish evolutionary links (if they exist) between different classes of isolated NSs, and to understand better the general evolution of these sources."
    },
    "0910/0910.0640_arXiv.txt": {
        "abstract": "Using a parametric approach, we determine the configuration of super-AGB stars at the explosion as a function of the initial mass and metallicity, in order to verify if the EC-SN scenario involving a super-AGB star is compatible with the observations regarding SN2008ha and SN2008S. The results show that both the SNe can be explained in terms of EC-SNe from super-AGB progenitors having a different configuration at the collapse. The impact of these results on the interpretation of other sub-luminous SNe is also discussed. ",
        "introduction": "\\label{sec:intro}  It is widely accepted that there are two main explosion mechanisms leading to supernova (SN) events \\citep[e.g.][]{woosley86,hille00}: the thermonuclear runaway in white dwarfs approaching the Chandrasekhar mass and the core collapse of stars with initial mass $\\gtrsim 8$\\Msun (CC-SNe). From an observational point of view, the former mechanism originates the relatively homogeneous type Ia SNe. The latter produces a huge variety of displays (different energetics and amounts of ejected material, reflecting on heterogeneous light curves and spectra evolutions), which are thought to be linked to the possible interaction of the CC-SN ejecta with circumstellar material (CSM) and the different configuration of the CC-SN progenitor at the time of the explosion \\citep*[e.g.][]{filip97,hamuy03,turatto03,turatto07}.  However the exact nature of CC-SN progenitors (initial mass, stellar structure and composition at the explosion) and the type of collapse (iron core collapse or neon-oxygen core collapse triggered by electron captures) are far from being well-established. There are still ambiguities that arise, on the theoretical side, from the uncertainties in modelling stellar evolution and the explosion mechanism \\citep[e.g.][]{woosley02,heger03} and, on the observational side, from the sparse direct detections of progenitor stars and a controversial classification of some events \\citep[e.g.][]{smartt08,turatto07}.  A number of faint transients have been recently discovered whose nature is still ambiguous and extensively debated. In particular SN~2008S received a SN designation by \\citet[][]{stan08}, but \\citet[][]{steele08} considered it as a SN ``impostor'', and \\citet[][]{smith08} as the exotic eruption of a luminous blue variable (LBV) object with a relatively low-mass, highly obscured progenitor ($\\lesssim 15$\\Msun). An eruptive origin was invoked also for two other similar transients \\citep[M 85 OT2006-1 and NGC 300 OT2008-1;][]{kulk07,Berger09,bond09}. However, work based on multi-wavelength follow-up of the transients and mid-IR images analysis of the pre-explosion environments not only did not rule out a SN origin \\citep[][]{pastorello07,Prieto08} but even suggested that they may be CC-SNe triggered by electron-capture reactions (so-called EC-SNe) involving a super-asymptotic giant branch (super-AGB) star \\citep[][]{thompson08}. The long-term multiwavelength monitoring of the SN2008S and new comparisons with the two aforementioned transients seem to support the EC-SN interpretation \\citep[][B09 hereafter]{bot09}. In particular B09 favor a scenario where the SN2008S progenitor is a super-AGB star embedded in an optically thick circumstellar shell. This conclusion is based on (1) the fact that the pre-explosion luminosity of the progenitor is in plausible agreement with the super-AGB models, (2) the capability of super-AGB stars to form thick circumstellar shells through mass-loss during the thermal pulses phase, (3) the similarity in the total radiated energy of SN2008S with that of other faint SNe, (4) the moderate velocities ($\\sim 3000$\\kms) of the ejecta, and (5) a low but significant mass ($0.0014 \\pm 0.0003$\\Msun) of ejected $^{56}$Ni.  An EC-SN explanation has been suggested also for SN2008ha \\citep[][V09 hereafter]{val09}, although an iron CC-SN involving a massive star (initial mass $\\gtrsim 25-30\\,M_{\\odot}$) plus fallback onto the collapsed remnant can not be ruled out. At first this object was included among the SN2002cx-like variety of peculiar type Ia SNe \\citep[][]{Li03,jha06,phillips07}, but V09 reviewed the thermonuclear scenario on the basis of the photometric and spectroscopic similarity to low-luminosity CC-SNe, concluding that all SN 2002cx-like objects could be indeed faint, stripped-envelope CC-SNe and that SN2008ha represents the faint tail in the luminosity distribution of this SN family. However, \\citet{foley09} did not definitively rule out the thermonuclear origin of the SN2002cx-like objects, and proposed to explain the peculiarity of SN2008ha in terms of an ``accretion-induced'' collapse \\citep[so-called AIC mechanism; see][for details]{metzger09}.  So far no clear picture has emerged and different scenarios may explain the aforementioned faint transients, especially because a detailed scrutiny of the super-AGB progenitors configuration at the explosion, which is crucial for a comparison with SN observables, is still missing. In the light of the most recent super-AGB stars models (e.g. \\citealt{sp06}, SP06 hereafter; \\citealt{thesis}, P06 hereafter; \\citealt{s07}, S07 hereafter; \\citealt{Poel08}), in this Letter we discuss in detail the possible link between EC-SNe from super-AGB progenitors and these transients using a parametric approach. After a brief sketch of the super-AGB stars evolution, we determine their configuration at the explosion as a function of the initial mass and metallicity from the most recent grids of super-AGB stellar models, and then we investigate if such a configuration is compatible with the observations. ",
        "conclusions": "The two events SN2008ha and SN2008S find a reasonable interpretation in the aforementioned scenario, and the progenitor mass to be associated with these SNe can be determined, considering the best ``global'' matching between the features of the super-AGBs models and the observed properties.  Assuming an initial metallicity $Z$ for the progenitors from $\\sim 0.008$ to $\\sim 0.02$ (see e.g. V09; B09; \\citealt{foley09}, for details on the metallicity determination), one obtains that a super-AGB star with initial mass slightly above $M_{N}$ has M$_{core}^{postCB} \\lesssim 1.25$-$1.26$\\Msun\\, (cf. second row in the sets of models in Tab.~\\ref{tabl1}), while a super-AGB star with initial mass $\\sim (M_{N}$+$0.5)$\\Msun\\, has M$_{core}^{postCB} \\sim 1.34$-$1.35$\\Msun\\, (cf. the row before the last in the sets of models in Tab.~\\ref{tabl1}). As a consequence the time $\\Delta t_{CH}$ necessary to the H-free core to reach $M_{CH}$ is $\\gtrsim 2.2$-$2.5\\cdot10^{5}\\,yr$ in the former case, and $\\sim 5\\cdot10^{4}\\,yr$ in the latter one. This difference in the time elapsing between the beginning of TP-SAGB phase and the EC-SN event in the two cases, reflects on the configuration at the collapse. The super-AGB star with initial mass slightly above $M_{N}$ has time to expel almost all the envelope and, consequently, gives rise to a faint stripped-envelope SN characterised by a non H-rich\\footnote{ We do not have accurate quantitative informations about the chemical composition of the ejecta (except for the $^{56}$Ni) to be compared to observations of SN2008ha, due to uncertainties of both observational and theoretical nature. However, we speculate that the composition could be non H-rich. In fact, for this model the ejecta is composed by the H-free stellar layer between the mass cut and $M_{CH}$ (representing $\\sim$ 5-15\\% by mass of all the ejected mass) and by the remaining envelope mass at the explosion, whose ``initial'' H-rich composition can be deeply altered by the second dredge-up phenomenon, the so-called Hot Botton Burning, and the third dredge-up episodes.} ejecta of $\\lesssim 0.2$-$0.3$\\Msun\\, with no signatures of prompt CSM interaction, in agreement with the observations of SN2008ha (M$_{ej}$ in the range $0.1$-$0.5$\\Msun\\,, e.g. V09; \\citealt{foley09}). Assuming an average wind velocity of $10$\\kms\\,, $90$\\% of the total expelled mass can be at a radial distance $\\gtrsim 5 \\cdot 10^{^4}A.U.$ when the EC-SN event takes place. The mean density of the CSM is expected to be $\\lesssim5$cm$^{-3}$ (this value is likely to be even lower due to a decreased mass loss rate near the end of the TP-SAGB phase when the mass of the envelope is significantly reduced), that could be sufficiently low not to give rise to significant interaction. The relatively low X-ray emission \\citep[L$_{X} < 5 \\cdot 10^{39}$ erg\\,s$^{-1}$;][]{foley09} seems to support this idea, because the CSM can be an efficient X-ray radiator for much higher density \\citep[$\\sim 10^5$-$10^7$cm$^{-3}$;][]{Cheva01}.  On the contrary, the super-AGB star with initial mass $\\sim (M_{N}$+$0.5)$\\Msun\\, loses $\\sim 1.6$-$1.8$\\Msun\\, in $\\sim 5\\cdot10^{4}\\,yr$ --- consistently with the inferred duration of the so-called dust-enshrouded phase for SN2008S \\citep[upper limit equal to $\\sim 6\\cdot 10^{4}\\,yr$;][]{thompson08} --- and, besides maintaining a massive ($\\sim 7.4$\\Msun) envelope at the collapse, could be embedded within a thicker circumstellar envelope (mean density $\\sim90$cm$^{-3}$). Observations of SN2008S \\citepalias[][]{bot09} indicate the formation of a detached shell with an inner radius of $\\sim 90\\, A.U.$ and outer radius of $\\sim 450\\,A.U.$ having $\\sim 0.08$\\Msun\\, of gas (van Loon, private communication). We could reproduce such a shell increasing the mass loss rate by $\\sim 15$ times above the average value for a relatively short period of $\\sim 150\\, yr$ as a consequence of a He-shell flash episode \\citep[see][for details]{Matt07}. In addition, we find that $\\sim 95$\\% of the total expelled mass in the CSM is beyond the aforementioned detached shell, and these findings are in agreement with the presence of dust at a radial distance greater than $\\sim 2 \\cdot 10^{^3}A.U.$, as inferred from the MIR emission of SN2008S \\citepalias[][]{bot09}.  Thus the current understanding of super-AGB stellar evolution is quantitatively consistent with the available data on these two recent faint transients, that may be explained in terms of EC-SNe from super-AGB progenitors having a different configuration at the collapse, without resorting to ``exotic'' scenarios that are not free from uncertainties. As for the ``special'' eruption of LBV of relatively low mass proposed to explain the features of the SN2008S \\citep[][]{smith08}, in addition to the problems for reconciling the ejecta velocity $\\lesssim 3000$\\kms\\, with a stellar eruption (B09), it is difficult to explain the fact that the slope of the late-time light curve of SN2008S \\citep[but also that of the similar event NGC300-OT;][]{bond09} is surprisingly similar to that expected in a SN explosion when the main mechanism powering the SN luminosity is the radioactive decay of $^{56}$Co into $^{56}$Fe. As for the AIC mechanism invoked for the SN2008ha, the main problem concerns the high velocity ($\\sim 0.1$-$0.2 c$) not observed in the ejecta and the impossibility to synthesise the observed intermediate-mass nuclei, that are predicted by the ``standard'' (involving a single degenerate binary system) AIC model. The so-called ``enshrouded'' AIC model involving the merging of two WDs in a binary system \\citep[][]{metzger09} might be somewhat less problematic. However the ejecta velocity, the amount of $^{56}$Ni and the production of intermediate-mass elements are still quantitatively poorly defined, and the role of the possible interaction between the disk wind and the outgoing SN shock has to be explored.  The weakness of the explosion and small amount of $^{56}$Ni synthesized make EC-SNe an obvious explanation for low-luminosity core-collapse events with unusual properties that are related to the pre-explosion mass loss episodes of their super-AGB progenitors and/or to the possible ensuing ejecta/CSM interaction. However, it has been suggested that the EC channel may also account for the properties of some relatively ``normal'' type II SNe \\citep[e.g.][]{CU00,kitaura06}, characterised by low luminosity, small amount of ejected $^{56}$Ni, extended plateaus (implying envelope mass of several \\Msun) and slow expansion velocities \\citep[e.g.][]{pastorello09}. To date, only for two objects of this class \\citep[SN2005cs and SN2008bk;][]{mau06,li06,matt08} clear evidence has been found for low mass progenitors on pre-explosion images, but the fact that they are super-AGBs is strongly questionned \\citep*[e.g.][]{eldridge07}. Thus, it remains to be seen what fraction (if any) of low luminosity type II SNe are EC-SNe and what other, instead, are more usual iron CC-SNe that experience less energetic than normal explosions \\citep[as, for example, if some of them are sufficiently massive to undergo fallback onto the collapsed remnant; see e.g.][]{Zampieri03}.  The wide variety of displays expected for EC-SNe may be of interest also in understanding the two unusual events, SN2005E and SN2005cz \\citep[][]{kawa09,perets09}. Indeed this scenario can account for many of the observed characteristic of both SNe (namely low explosion energy, very low ejected mass and ejection of small amount of $^{56}$Ni), but the possibility to reproduce all the observed properties (as the spectroscopic features and, in particular, the alleged Ca-richness) deserves further investigation.  We are aware that large uncertainties of both theoretical and observational nature are still present on the EC-SN mechanism in super-AGB stars. Nevertheless we believe that the scenario herein proposed is promising for understanding an increasing number of underenergetic and unusual SNe. Only a combined effort will solve the issue. On one side we need more accurate observational constraints about the production of intermediate-mass nuclei (specifically C, O and all the $\\alpha$-elements in general) in low luminosity SN events. On the other side more refined future studies on the super-AGB stellar evolution fully describing the TP-SAGB phase, and 3-D simulations for examining in detail the nucleosynthesis processes in EC-SNe are desirable."
    },
    "0910/0910.5253_arXiv.txt": {
        "abstract": "Using astrometric VLBI observations, we have determined the parallax of the black hole X-ray binary V404 Cyg to be $0.418\\pm0.024$\\,mas, corresponding to a distance of $2.39\\pm0.14$\\,kpc, significantly lower than the previously accepted value.  This model-independent estimate is the most accurate distance to a Galactic stellar-mass black hole measured to date.  With this new distance, we confirm that the source was not super-Eddington during its 1989 outburst.  The fitted distance and proper motion imply that the black hole in this system likely formed in a supernova, with the peculiar velocity being consistent with a recoil (Blaauw) kick.  The size of the quiescent jets inferred to exist in this system is $<1.4$\\,AU at 22\\,GHz.  Astrometric observations of a larger sample of such systems would provide useful insights into the formation and properties of accreting stellar-mass black holes. ",
        "introduction": "Stellar-mass black holes in accreting X-ray binary systems provide unique probes of the physical properties and formation mechanism of black holes, as well as the physics of accretion and associated outflow.  However, our understanding of black hole systems is hampered by the fact that distances to X-ray binaries are relatively poorly known, typically uncertain by 50\\% or more \\citep{Jon04}.  Standard spectroscopic estimates are handicapped by the distortions induced in the secondary by its compact partner, contamination by emission from the accretion disc \\citep{Rey08}, and uncertainties in estimates of interstellar absorption.  Since these distance uncertainties are systematic rather than random, they cannot be reduced by making more observations.  According to \\citet{Jon04}, these systematic uncertainties can artificially enlarge the difference between the quiescent X-ray luminosities of black hole and neutron star systems. If true this would undermine the claimed evidence for an event horizon \\citep{Nar97,Men99,Gar01}. Accurate distances are also required to interpret the measured proper motions of black hole X-ray binaries, by converting a measured proper motion into a physical speed, from which we can derive the peculiar velocity of the source \\citep[e.g.,][]{Mir01,Mir02,Mir03,Dha07,Mil09}. This can be used to place constraints on any velocity kick the black hole might have received in its natal supernova, and hence on the formation mechanism of the black hole \\citep{Bra95,Wil05,Fra09}.  The only direct, model-independent method of measuring distances is via trigonometric parallax, which has hitherto been impossible for black hole X-ray binary systems, since they lie at distances of several kpc, and hence require sub-milliarcsecond astrometry to measure their parallaxes.  Very Long Baseline Interferometry (VLBI) at radio wavelengths is currently the only technique available for such high-precision astrometric measurements.  However, black hole X-ray binaries spend the majority of their time in a faint quiescent state, where they are often not detected in the radio band at current sensitivities.  They become bright enough to be detected during outbursts, but such outbursts are typically not sufficiently long or frequent and do not sample the Earth's orbit well enough to make a parallax measurement feasible.  During outbursts, the radio emission is often resolved at the high angular resolutions required for precision astrometry, making it difficult to determine the true location of the binary system from epoch to epoch.  Lastly, many black hole X-ray binaries are located in or close to the Galactic Plane, such that the scatter-broadening along the line of sight makes high-precision astrometry impossible in the centimeter waveband where current VLBI arrays have the highest sensitivity.  Here we present High Sensitivity Array (HSA) observations of V404 Cyg, the most luminous known black hole X-ray binary in quiescence.  With a mass function of $6.08\\pm0.06M_{\\odot}$ \\citep{Cas94}, the compact object is a dynamically-confirmed black hole, accreting matter from a K0 subgiant companion star \\citep{Cas93}.  The system has a quiescent radio flux density of 0.3\\,mJy, with a flat spectrum indicative of a self-absorbed compact jet \\citep{Gal05}, although the jets are not resolved at the angular resolution of the HSA \\citep{Mil08}.  The persistent, unresolved radio emission makes this source a good target for astrometric observations to measure its parallax. ",
        "conclusions": "We have used HSA observations to measure the most accurate distance to a black hole to date, via the method of trigonometric parallax.  The distance of V404 Cyg is $2.39\\pm0.14$\\,kpc, significantly lower than the previous best estimate of the source distance, and demonstrating that uncertainties in the interstellar extinction can introduce significant errors into distance estimates.  With the new distance, we can determine that V404 Cyg did not exceed its Eddington luminosity during the 1989 outburst, and that it is likely to have formed in a natal supernova, with the peculiar velocity being consistent with recoil from the ejected mass, and any asymmetric kick being small.  We have further constrained the size of the quiescent jets in the source to $<1.4$\\,AU.  These observations demonstrate the feasibility of measuring parallax distances to black hole X-ray binaries using VLBI techniques, which could, with a larger sample of systems, firm up evidence for black hole event horizons, and provide new insights into black hole spins, the mechanism of black hole formation, and the nature of ULXs."
    },
    "0910/0910.3196_arXiv.txt": {
        "abstract": "In this paper we extend a previous work~\\citep{vielva09} where we presented a method based on the N-point probability distribution (pdf) to study the Gaussianity of the cosmic microwave background (CMB). We explore a local non-linear perturbative model up to third order as a general characterization of the CMB anisotropies. We focus our analysis in large scale anisotropies ($\\theta > 1^\\circ$). At these angular scales (the Sachs-Wolfe regime), the non-Gaussian description proposed in this work defaults (under certain conditions) to an approximated local form of the weak non-linear coupling inflationary model. In particular, the \\emph{quadratic} and \\emph{cubic} terms are governed by the non-linear coupling parameters f$_\\nl$ and g$_\\nl$, respectively. The extension proposed in this paper allows us to directly constrain these non-linear parameters. Applying the proposed methodology to WMAP 5-yr data, we obtain $-5.6\\times10^5 < {\\rm g}_{\\nl} < 6.4\\times10^5$, at 95\\% CL. This result is in agreement with previous findings obtained for equivalent non-Gaussian models and with different non-Gaussian estimators, although this is the first direct constrain on g$_\\nl$ from CMB data. A model selection test is performed, indicating that a Gaussian model (i.e. f$_\\nl \\equiv 0$ and g$_\\nl \\equiv 0$  ) is preferred to the non-Gaussian scenario. When comparing different non-Gaussian models, we observe that a pure f$_\\nl$ model (i.e. g$_\\nl \\equiv 0$) is the most favoured case, and that a pure g$_\\nl$ model (i.e. f$_\\nl \\equiv 0$) is more likely than a general non-Gaussian scenario (i.e. f$_\\nl \\neq 0$ and g$_\\nl \\neq 0$). Finally, we have analyzed the WMAP data in two independent hemispheres, in particular the ones defined by the dipolar pattern found by~\\cite{hoftuft09}. We show that, whereas the g$_\\nl$ value is compatible between both hemispheres, it is not the case for f$_\\nl$ (with a p-value $\\approx 0.04$). However, if, as suggested by~\\cite{hoftuft09}, anisotropy of the data is assumed, the distance between the likelihood distributions for each hemisphere is larger than expected from Gaussian and anisotropic simulations, not only for f$_\\nl$, but also for g$_nl$ (with a p-value of $\\approx 0.001$ in the case of this latter parameter). This result is an extra evidence for the CMB asymmetries previously reported in WMAP data. ",
        "introduction": "\\label{intro}   Current observations of the Cosmic Microwave Background (CMB) temperature and polarization fluctuations, in addition to other astronomical data sets~\\citep[e.g.][see \\citealt{barreiro09} for a recent review]{komatsu09,gupta09}, provide an overall picture for the origin, evolution, and matter and energy content of the universe, which is usually referred to as the \\textit{standard cosmological model}. In this context, we believe the universe to be highly homogeneous and isotropic, in expansion, well described by a Friedmman-Robertson-Walker metric and with a trivial topology. The space geometry is very close to flat, and it is filled with cold dark matter (CDM) and dark energy (in the form of a cosmological constant, $\\Lambda$), in addition to baryonic matter and electromagnetic radiation. Large scale structure (LSS) is assumed to be formed by the gravitational collapse of an initially smooth distribution of adiabatic matter fluctuations, which were seeded by initial Gaussian quantum fluctuations generated in a very early inflationary stage of the universe evolution.   It is interesting to mention that, besides the success of current high quality CMB data~\\citep[in particular the data provided by the WMAP satellite][]{hinshaw09} in constraining the cosmological parameters with good accuracy and in showing the high degree of homogeneity and isotropy of the Universe~\\citep[as predicted by the standard inflation scenario, see for instance][]{liddle00}, it has been, precisely, through the analysis of this very same data, that the CMB community has been allowed to probe fundamental principles and assumptions of the \\emph{standard cosmological model}. Most notably, the application of sophisticated statistical analysis to CMB data might help us to understand whether the temperature fluctuations of the primordial radiation are compatible with the fundamental isotropic and Gaussian standard predictions from the \\emph{inflationary phase}.  Indeed, the interest of the cosmology community in this field has experienced a significant growth, since several analyses of the WMAP data have reported some hints for departure from isotropy and Gaussianity of the CMB temperature distribution. The literature on the subject is very large, and is still growing, which makes really difficult to provide a complete and updated list of publications. We refer to our previous work~\\citep{vielva09} for an (almost) complete review.  Among the previously mentioned analyses, those related to the study of non-standard inflationary models have attracted a larger attention. For instance, the non-linear coupling parameter f$_{\\nl}$ that describes the non-linear evolution of the inflationary potential \\citep[see e.g.][and references therein]{bartolo04} has been constrained by several groups, from the analysis of the WMAP data: using the angular bispectrum \\citep{komatsu03,creminelli07,spergel07,komatsu09,yadav08,smith09}; applying the Minkowski functionals \\citep{komatsu03,spergel07,gott07,hikage08,komatsu09}; using different statistics based on wavelets \\citep{mukherjee04,cabella05,curto09a,curto09b,pietrobon09,rudjord09}, and by exploring the N-pdf \\citep{vielva09}. Besides marignal detections of f$_{\\nl} > 0$~\\citep[with a probability of around 95\\%,][]{yadav08,rudjord09}, there is a general consensus on the WMAP compatibility with the predictions made by the standard inflationary scenario at least at 95\\% confidence level. The current best limits obtained from the CMB data are: $-4 < {\\rm f}_{\\nl} < 80$ at 95\\% CL by~\\cite{smith09}. In addition, very recently, promising constraints coming out from the analysis of LSS have been reported: $-29 < {\\rm f}_{\\nl} < 70$ at 95\\% CL~\\citep{slozar08}.  The aim of this paper is to extend our previous work~\\citep{vielva09}, where the full N-pdf of a non-Gaussian model  that describes the CMB anisotropies as a local (pixel-by-pixel) non-linear expansion of the temperature fluctuations (up to second order) was derived. For this model ---that, at large scales, can be considered as an approximation to the weak non-linear coupling scenario--- we are able to build the exact likelihood on pixel space. Working in pixel space allows one to include easily non-ideal observational conditions, like incomplete sky coverage and anisotropic noise. The extension of the present work is to account for higher order moments in the expansion, in particular, we are able to directly obtain constraints on g$_\\nl$, that is the coupling parameter governing the cubic term of the weak non-linear expansion.  As far as we are aware, direct constraints on g$_\\nl$ have been made available only very recently \\citep{desjacques09} and are obtained from LSS analyses: $-3.5\\times 10^{5} < {\\rm g}_{\\nl} < 8.2\\times 10^{5}$ at 95\\% CL. This constraint was obtained for the specific case in which the coupling parameter governing the quadratic term (f$_\\nl$) is negligible~\\citep[i.e. f$_\\nl \\equiv 0$, which is the situation required for some curvaton inflationary models, e.g.][]{sasaki06,enqvist08,huang09}. We present in this work the first direct measurement of g$_\\nl$ obtained from CMB data. In addition to study the particular case of f$_\\nl \\equiv 0$, we also consider a more general case in which a joint estimation of f$_\\nl$ and g$_\\nl$ is performed. Finally, and justified by recent findings~\\citep[e.g.,][]{hoftuft09}, we compute the N-pdf in two different hemispheres, and derive from it constraints on f$_\\nl$ and g$_\\nl$ for this hemispherical division of the celestial sphere.   The paper is organized as follows. In Section~\\ref{sec:model} we describe the physical model based on the local expansion of the CMB fluctuations and derive the full posterior probability, recalling how it defaults to the case already addressed in \\cite{vielva09}. In Section~\\ref{sec:simulations} we check the methodology against WMAP-like simulations. Results on WMAP 5-year data are presented in Section~\\ref{sec:wmap}. Conclusions are given in Section~\\ref{sec:final}. Finally, in Appendix~\\ref{sec:app}, we provide a detailed computation of the full N-pdf. ",
        "conclusions": "\\label{sec:final}  We have presented an extension of our previous work~\\citep{vielva09}, by defining a parametric non-Gaussian model for the CMB temperature fluctuations. The non-Gaussian model is a local perturbation of the standard CMB Gaussian field, which (under certain circumstances) is related to an approximative form of the weak non-linear coupling inflationary model at scales larger than the horizon scale at the recombination time \\citep[i.e., above the degree scale, see for instance][]{komatsu01,liguori03}. For this model, we are able to build the posterior probability of the data given the non-linear parameters f$_\\nl$ and g$_\\nl$. From these pdfs, optimal maximum likelihood estimators of these parameters can be obtained.  We have verified with WMAP-like simulations that the maximum likelihood estimation of the \\emph{quadratic} non-linear coupling parameters ($\\hat{\\rm f}_\\nl$) is unbiased, at least for a reasonable range of values, even when non-Gaussian simulations also account for a \\emph{cubic} term. In particular, we found that for simulated non-Gaussian coefficients such as $|\\bar{\\rm f}_\\nl| \\lesssim 400$ and $|\\bar{\\rm g}_\\nl| \\lesssim 5\\times 10^{5}$, the estimation of f$_\\nl$ is accurate and efficient. However, when trying to study the case in which only a pure \\emph{cubic} model is addressed, the situation is different. In particular, the simulated \\emph{quadratic} term has an important impact on the estimation of g$_\\nl$. For instance, if a pure \\emph{quadratic} non-Gaussian model is simulated with a value of $\\bar{\\rm f}_\\nl \\equiv 400$, a value of $\\hat{\\rm g}_\\nl \\approx 2.5\\times 10^5$ is wrongly estimated. This results indicates, obviously, that the \\emph{quadratic} term is more important than the \\emph{cubic} one in the expansion of the local non-Gaussian model, and, therefore, that not accounting properly for the former might have a dramatic impact on the latter. Contrarily, the opposite situation is much more unlikely. Finally, we have investigated the joint estimation of the f$_\\nl$ and g$_\\nl$ parameters. In this case we find that for a similar regime as the one mentioned above (i.e., $|\\bar{\\rm f}_\\nl| \\lesssim 400$ and $|\\bar{\\rm g}_\\nl| \\lesssim 5\\times 10^{5}$), an accurate and efficient estimation of the non-linear coupling parameters is obtained. However, for larger values of these coefficients, we find that the parameter estimation is highly biased, favouring a region of the parameter space of larger values for both coefficients.  We have addressed, afterwards, the analysis of the WMAP 5-year data. We have consider two different analyses. First, we have investigated the case of an all-sky analysis (except for the Galactic area not allowed by the WMAP KQ75 mask). Second, and motivated by previous findings, we have performed a separated analysis of two hemispheres. In particular, the hemispherical division associated to the dipolar pattern found by~\\cite{hoftuft09} was considered.  Regarding the all-sky analysis, we find, for the case in which a pure \\emph{quadratic} model is investigated, the same result that we already found in our previous work~\\citep{vielva09}. In particular, we determine that $-154 < {\\rm f}_\\nl < 94$ at 95\\%. Equivalently, for the case of a pure \\emph{cubic} non-Gaussian model we establish $-5.6\\times 10^5 < {\\rm g}_\\nl < 6.4\\times 10^5$ at 95\\%. This is in agreement  with a recent work by~\\cite{desjacques09}, where an analysis of LSS data was performed. The result that we provide in this paper is, as far as we know, the first direct constraint on g$_\\nl$ from CMB data alone. Finally, we have also investigated the case of a joint estimation of the \\emph{quadratic} and \\emph{cubic} non-linear coupling parameters. In this case, the constraints obtained on f$_\\nl$ and g$_\\nl$ are virtually the same as the ones already reported for the independent analyses.  We have performed a model selection to evaluate which of the four hypotheses (i.e. Gaussianity, a pure \\emph{quadratic} model, a pure \\emph{cubic} model and a general non-Gaussian scenario) is more likely. Since a well motivated prior for the non-linear coupling parameters is lacking, we have used asymptotic model selection criteria (like AIC and BIC), instead of more powerful Bayesian approaches, like the Bayesian Evidence or the posterior ratio test. Both methodologies (AIC and BIC) indicates that the Gaussian hypothesis is more likely than any of the non-Gaussian models. We also found that, among the non-Gaussian scenarios, the one with a pure \\emph{quadratic} model is the most favoured one, whereas the general one (i.e. f$_\\nl \\ne 0$ and g$_\\nl \\ne 0$) is the most unlikely.   The analysis of the WMAP data in two hemispheres revealed that, whereas both halves of the sky present similar constraints on the g$_\\nl$ parameter (and, in both cases, not indicating any significant incompatibility with the zero value), the analysis of a pure \\emph{quadratic} scenario showed a clear asymmetry. First, the f$_\\nl$ value in the hemisphere whose pole is closer to the Southern Ecliptic Pole, is $\\hat{\\rm f}_\\nl = -164 \\pm 62$. Which implies that f$_\\nl < 0$ at 96\\%. Even more, the distance between both likelihoods (given in terms of the KLD) presents a p-value of $\\approx 0.04$.  We have also analyzed the WMAP data after by considering different correlation properties in each hemisphere (according to the dipolar modulation described by~\\cite{hoftuft09}). We have tested that the behaviour found for the f$_\\nl$ is practically the same as before, and that, in addition, an extra anomaly appears associated with the g$_\\nl$ parameter. In particular, the distance between both likelihoods is anomalously large as well (it corresponds to a p-value of $\\lesssim 0.002$. Further test was performed after correcting WMAP data from the dipolar modulation. In this case the asymmetries in the maximum-likelihood estimations of both non-linear coupling parameters remain unaltered. Hence, these results indicate that, as it has been previously reported in other works, there are evidences of anisotropy in the WMAP data, reflected as an asymmetry between two opposite hemispheres. Such anomaly is related to a different distribution for a non-linear coupling parameter related to the \\emph{quadratic} term. However, a correction in terms of a dipolar modulation as the one proposed by~\\cite{hoftuft09}, seems not to be sufficient to account for this anomaly related to the likelihood distribution of the f$_\\nl$ parameter."
    },
    "0910/0910.4580_arXiv.txt": {
        "abstract": "The evolution of the expansion rate of the Universe results in a drift in the redshift of distant sources over time. A measurement of this drift would provide us with a direct probe of expansion history. The Lyman $\\alpha$ (\\lya) forest has been recognized as the best candidate for this experiment, but the signal would be weak and it will take next generation large telescopes coupled with ultra-stable high resolution spectrographs to reach the cm~s$^{-1}$ resolution required. One source of noise that has not yet been assessed is the transverse motion of \\lya~absorbers, which varies the gravitational potential in the line of sight and subsequently shifts the positions of background absorption lines. We examine the relationship between the pure cosmic signal and the observed redshift drift in the presence of moving \\lya~clouds, particularly the collapsed structures associated with Lyman limit systems (LLSs) and damped Lyman $\\alpha$ systems (DLAs). Surprisingly, the peculiar velocities and peculiar accelerations {\\it both} enter the expression, although the acceleration term stands alone as an absolute error, whilst the velocity term appears as a fractional noise component. An estimate of the magnitude of the noise reassures us that the motion of the \\lya~absorbers will not pose a threat to the detection of the signal. ",
        "introduction": "\\label{intro}  The large-scale homogeneity and isotropy of the Universe is encompassed in Friedmann-Lema\\^{i}tre-Robertson-Walker (FLRW) geometry (See \\citet*{H06} for further details). One of the key aspects of this cosmology is the evolution of the scale factor, $R(t)$; modern cosmology aims to characterize this evolution, distinguish between competing cosmological models and predict the fate of the Universe. Observationally, this has been a rather difficult science, with so-called standard candles and standard rulers playing leading roles in indirect measurements. \\citet{S62} recognized that a direct detection of the evolution of the expansion rate was, in theory, possible; just as the cosmological redshifts of extra-galactic objects in all directions are a consequence of isotropic expansion, a positive or negative drift in the redshift {\\it of a single object} over time would indicate acceleration or deceleration with respect to the observer. The redshift drift is very small, but as we head towards an era of cm~s$^{-1}$ resolution spectroscopy, the prospect of detection becomes very real. In addition to identifying the appropriate instrumentation and spectroscopic techniques required, it is crucial to account for any systematic biases and consider if any source of noise may hinder the detection of the cosmic signal. Previous studies have covered some of this ground but the potential for the transverse motion of Lyman~$\\alpha$ (\\lya) clouds to hide this signal remains yet unexplored. In the present paper we aim to address this issue and estimate the magnitude of the bias or noise that is introduced.  In Section~\\ref{background} we review the current cosmological paradigm along with relevant observational tests. We re-derive the expression for the cosmic signal being sought in Section~\\ref{cosmicsignal}. In Section~\\ref{movinglenseffect} we introduce the transverse moving lens effect. We derive the expression for the observed redshift drift in the presence of \\lya~clouds in Section~\\ref{zdrift} and discuss the results in Section~\\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  The redshift drift experiment has been recognized as a direct test of the cosmic expansion, and carries with it the potential to independently constrain cosmological parameters and distinguish between evolving dark energy models. The proposed next generation of ELTs and ultra-stable high resolution spectrographs make the detection of the cosmic signal possible in the near future. However, as the signal is weak, it is necessary to account for all possible systematic biases and sources of noise. One source of noise that has, until now, remained unexplored is the variation in the gravitational field due to the transverse motion of \\lya~absorbers. In this paper, we have examined this effect, and characterised and quantified the noise introduced to the signal. To summarise, we find that the peculiar velocities and accelerations of the \\lya~clouds both enter the expression relating the observed redshift to the cosmic signal as noise terms; the former introduces a fractional offset, the latter, absolute. The evolution of the gravitational potential wells in the line of sight due to transverse motion of the \\lya~absorbers will not, in fact, introduce significant noise or bias into the detection of the cosmic signal.  We consider now whether this is the final word on evolving gravitational fields. \\lya~absorbers were the most obvious source of foreground interference; they, by definition, intersect the line of sight. Galactic haloes at impact parameters of 1--20~kpc are thought to produce the rarer absorption effects, such as LLSs and DLAs; these have been the extreme cases studied in this work. Estimates of the frequency and redshift distribution of these objects are hindered by reddening, which can dim quasars and exclude them from flux-limited samples, although the magnification bias can counteract this effect. Both effects can be expected, particularly if the absorbers are massive galaxy haloes at low impact parameters. The resulting evolution of the gravitational field within the filaments and sheets associated with the low-density \\lya~forest is difficult to study analytically. The filaments feed into more massive structures, such as galaxy clusters, with absorption taking place at impact parameters of $\\sim$1~Mpc. The transverse peculiar motion of such objects will have a similar `moving-lens' effect, the rarity of which has not been assessed. Their presence may not be noticed in the quasar spectra, though perhaps in direct imaging (at least at low redshift). Further studies using dark matter simulations will probe the expected motions of all mass scales in the vicinity of the line of sight to distant quasars, allowing the determination of the full influence of the BG effect in future redshift drift experiments; this will form the basis of future contributions."
    },
    "0910/0910.4125_arXiv.txt": {
        "abstract": "We expect a detectable correlation between two seemingly unrelated quantities: the four point function of the cosmic microwave background (CMB) and the amplitude of flux decrements in quasar (QSO) spectra. The amplitude of CMB convergence in a given direction measures the projected surface density of matter. Measurements of QSO flux decrements trace the small-scale distribution of gas along a given line-of-sight.  While the cross-correlation between these two measurements is small for a single line-of-sight, upcoming large surveys should enable its detection.  This paper presents analytical estimates for the signal to noise (S/N) for measurements of the cross-correlation between the flux decrement and the convergence, $\\langle\\delta \\flux \\kappa\\rangle$, and for measurements of the cross-correlation between the variance in  flux decrement and the convergence, $\\langle(\\delta \\flux)^2 \\kappa\\rangle$.  For the ongoing BOSS (SDSS III) and Planck surveys, we estimate an S/N of 30 and 9.6 for these two correlations.  For the proposed BigBOSS and ACTPOL surveys, we estimate an S/N of 130 and 50 respectively.  Since $\\langle(\\delta \\flux)^2 \\kappa\\rangle \\propto \\sigma_8^4$, the amplitude of these cross-correlations can potentially be used to measure the amplitude of $\\sigma_8$ at $z \\sim 2$ to 2.5\\% with BOSS and Planck and even better with future data sets.  These measurements have the potential to test alternative theories for dark energy and to constrain the mass of the neutrino.  The large potential signal estimated in our analytical calculations motivate tests with non-linear hydrodynamical simulations and analyses of upcoming data sets. ",
        "introduction": "The confluence of high resolution Cosmic Microwave Background (CMB) experiments and  large-scale spectroscopic surveys in the near future is expected to sharpen our view of the Universe. Arcminute scale CMB experiments such as Planck \\cite{PLANCK}, the Atacama Cosmology Telescope \\cite{ACT,hincks/etal:2009}, the South Pole Telescope  \\cite{SPT,staniszewski/etal:2009}, QUIET  \\cite{QUIET} and PolarBeaR \\cite{POLAR}, will chart out the small scale anisotropies in the CMB.  This will shed new light on the primordial physics of inflation, as well as the astrophysics of the low redshift Universe through the signatures of the interactions of the CMB photons with large scale structure.   Spectroscopic surveys like BOSS \\cite{mcdonald05,seljak05} and BigBOSS \\cite{Schlegel:2009uw} will trace the large scale structure of neutral gas, probing  the distribution and dynamics of matter in the Universe.  While these two datasets will be rich on their own, they will also complement and constrain each other. An interesting avenue for using the two datasets would be to utilize the fact that the arcminute-scale secondary anisotropies in the CMB are signatures of the same large scale structure that is  traced by the spectroscopic surveys, and study them in cross-correlation with each other. In this paper, we present the analytic estimates for  one such cross correlation candidate - that between the gravitational lensing of the CMB and the flux fluctuations in the  \\Lya forest. \\par The gravitational lensing of the CMB, or CMB lensing in short,  is caused by the deflection of the CMB photons by the large scale structure potentials  \\citep[see][for a review]{lewis.challinor:2006}. On large scales, WMAP measurements imply that the primoridal CMB is well described as an isotropic Gaussian random field \\cite{Komatsu:2008hk}.  On small scales, lensing breaks this isotropy and introduces a specific form of non-Gaussianity. These properties of the lensed CMB sky can be used to construct estimators of the deflection field that lensed the CMB. Therefore, CMB lensing provides us with a way of reconstructing a line-of-sight (los) projected density field from zero redshift to the last scattering surface, with a broad geometrical weighting kernel that gets most of its contribution from the $z=1-4$ range \\cite{hu.okamoto:2002,hirata.seljak:2003,yoo.zaldarriaga:2008}. While CMB lensing is mainly sensitive to the geometry and large scale projected density fluctuations, the \\Lya forest, the absorption in quasar (QSO) spectra caused by intervening neutral hydrogen in the intergalactic medium, primarily traces the small-scale distribution of gas (and hence, also matter) along the line of sight.  A cross-correlation between these two effects gives us a unique way to study how small scale  fluctuations in the density field evolve on top of large scale over and under-densities, and how gas traces the underlying dark matter. This signal is therefore a useful tool to test to what extent the fluctuations in the \\Lya flux relate to the underlying dark matter. Once that relationship is understood, it can also become a powerful probe of the growth of structure on a wide range of scales. Since both massive neutrinos and dark energy alter the growth rate of structure at $z\\sim 2$, these measurements can probe their effects. This new cross-correlation signal, should also be compared with other existing cross-correlations between CMB and LSS that have already been observed and that are sensitive to different redshift regimes \\cite{peiris.spergel:2000, giannantonio08, hirata08, croft06, Xia:2009dr}.  In this work, we build an analytic framework based on simplifying assumptions to estimate the cross-correlation of the first two moments of the \\Lya flux fluctuation with the weak lensing convergence $\\kappa$, obtained from CMB lensing reconstruction, measured along the same line of sight. The finite resolution of the spectrogram limits the range of parallel $k$-modes probed by the absorption spectra and the finite resolution of the CMB experiments limits the range of perpendicular $k$-modes probed by the convergence measurements. These two effect break  the spherical symmetry of the $k$-space integration. However, we show that by resorting to a power series expansion it is still possible to obtain computationally efficient expressions for the evaluation of the signal.  We then investigate the detectability of the signal in upcoming CMB and LSS surveys, and the extent to which such a signal can be used as a probe of neutrino masses and early dark energy scenarios. A highlight of our results is that the estimated cross-correlation signal seems to have significant sensitivity to the normalization of the matter power spectrum $\\sigma_8$. Consistency with CMB measurements -- linking power spectrum normalization and the sum of the neutrino masses -- allows to use this cross-correlation to put additional constrain on the latter. \\par  The structure of the paper is as follows. In Section \\ref{sectI} we  introduce the two physical observables, the \\Lya flux and the CMB convergence (\\ref{sectIa}), the cross-correlation estimators (\\ref{sectIb}) and their variances (\\ref{variance}).  Our main result is presented in section \\ref{sn} where the signal-to-noise ratios are computed. Section \\ref{spectral} contains a spectral analysis of the observables that aims at finding the \\Lya wavenumbers that contribute most to such a signal. We focus on two cosmologically relevant applications in sections \\ref{neutrinos} and \\ref{ede}, for massive neutrinos and early dark energy models, respectively. We conclude with a discussion in section \\ref{discuss}. ",
        "conclusions": "\\label{discuss} This work presents a detailed investigation of the cross-correlation signals between transmitted \\Lya flux and the weak lensing convergence of the CMB along the same line-of-sight. One of the motivations behind this work is that the \\Lya forest has already been shown to be a powerful cosmological tool and novel ways of exploring and deepening the understanding of the flux/matter relation could significantly improve our knowledge of the high redshift universe. These correlators are able to provide astrophysical and cosmological information: since they are sensitive to both the flux/matter relation and the value of cosmological parameters, in principle they can be used to put constraints on both.  The correlators investigated in the present work have a clear physical meaning. The correlation of $\\delta\\flux$ with $\\kappa$ measures to what extent the fluctuations along the los mapped by the \\Lya forest contribute to the CMB convergence field. This correlation is dominated by long wavelength modes ($k \\lesssim 10^{-1}\\,\\hMpcI$) and as such is more sensitive to \\Lya forest continuum fitting procedures. The correlation of the flux variance $\\delta\\flux^2$ with $\\kappa $ measures to what extent the growth of short wavelength modes (mapped by the \\Lya flux) is enhanced or depressed by the fact that the latter are sitting in regions that are overdense or underdense on large scales. This interplay between short and long wavelength modes is well exemplified by the redshift dependence of the S/N ratio for $\\langle\\delta\\flux^2\\kappa\\rangle$: lowering the redshift increases the S/N ratio because while the variance of $\\langle\\delta\\flux^2\\kappa\\rangle$ is dominated by the independent growth of long and short wavelength modes, the value of $\\langle\\delta\\flux^2\\kappa\\rangle$ itself receives an extra contribution due to the fact that the growth of the short wavelength modes is enhanced by the presence (and independent growth) of the long wavelength modes. Furthermore, this correlator is sensitive to intermediate-to-small scales ($k\\gtrsim 10^{-2}\\, \\hMpcI$) and as such it should be less sensitive to \\Lya forest continuum fitting procedures.  To estimate the values of the correlators, their variance and their S/N ratio we rely on linear theory and simple approximations, such as the fluctuating Gunn-Peterson approximation at first order. Although the framework is simplified, the results are by no means obvious since different modes enter non-trivially in these quantities and in their signal-to-noise ratio. We estimate that such correlations may be detectable at a high significance level by Planck and the SDSS-III BOSS survey, experiments that are already collecting data. Moreover, our investigation of the modes of the \\Lya forest that contribute to $ \\langle\\delta\\flux^2\\kappa\\rangle$ shows that the low-resolution \\Lya spectra measured by SDSS-III (which is aimed at the measurement of BAO at $z=2-4$ \\cite{mcdonald.eisenstein:2007,slosar09}) should have enough resolution to yield a significant S/N.  The peculiar dependence of $\\langle\\delta\\flux^2\\kappa\\rangle$ on intermediate- to-short scales and its sensitivity to the value of the power spectrum normalization $\\sigma_8$ makes it a very useful cosmological tool to test all models characterized by variations of the power spectrum on such scales. In particular, we applied our estimates to evaluate the sensitivity of $\\langle\\delta\\flux^2\\kappa \\rangle$ to changes in $\\sigma_8$ due to variations in the sum of the neutrino masses and to show how promising this measurement could be in constraining the latter.  Finally, some caveats are in order. First, the code developed to estimate $\\langle \\delta\\flux^2\\kappa\\rangle$ and its variance is based on the tree-level perturbation theory results reported here. As such, the results shown do not take into account nonlinearities induced by gravitational collapse. The extension of the analytic results to take into account this aspect is actually quite straightforward, as it only requires the implementation of the so-called ``HyperExtended Perturbation Theory'' for the bispectrum \\cite{Scoccimarro:2000ee}. However, the implementation of such changes in a numerical code are less trivial, as the integrations over the power spectrum and over the comoving distance cannot be factored any longer. We have nonetheless reason to speculate that the nonlinearities induced by gravitational collapse will not dramatically change the picture outlined here. At the redshift range spanned by the \\Lya forest nonlinearities are normally mild and confined to short scales. Furthermore, as shown in Sec.~\\ref{spectral}, the S/N ratio for $\\langle\\delta \\flux^2\\kappa\\rangle$ dominated by modes with $k\\gtrsim 10^{-2} \\,\\hMpcI$, but all decades above $10^{-2} \\,\\hMpcI$ contribute in the same proportion to both the signal \\textit{and} its variance. It is therefore conceivable to filter out of the \\Lya spectra the shortest scales, which are the most affected by nonlinearities and still be able to retain a non-negligible S/N.  The second caveat pertains the estimate of the correlators' variance. It is in fact necessary to point out that to obtain such \\textit{estimates} Wick's theorem has been applied. Whether the use of Wick's theorem may or may not lead to an accurate result when considering the variance of $\\langle \\delta\\flux^{2} \\kappa\\rangle$ is debatable. On one hand it is possible to point out that the largest part of the signal arises at small separations, where the value of the correlator is dominated by its connected part. Analogously, it could be possible to argue that the use of Wick's theorem may lead to underestimating the correlators' variance. An exact evaluation of the variance of $\\langle \\delta\\flux^{2} \\kappa\\rangle$, however, requires the exact calculation of a six point function, that to the best of our knowledge has never been determined. On the other hand it is also possible to point out that the connected part of $\\langle\\delta_c\\delta_{c'}\\delta_q^2\\delta_{q'}^2\\rangle$ will be significantly non-zero only when the distances between the different points is small. As such, this term will give a non-zero contribution proportional to the length of the \\Lya spectrum, which should be subdominant with respect to the ones considered in section \\ref{variance}, that are proportional to the distance from the observer all the way to the last scattering surface.  The third caveat pertains the expansion of the expression for the flux, Eq.~(\\ref {eq:FGPA}). Despite the fact that the expansion carried out in Eq.~(\\ref{deltaF}) is correct on scales larger than about 1 $\\hMpc$, we point out here that the flux as expressed in Eq.~(\\ref{eq:FGPA}) is intrinsically a non-linear function of the overdensity field. It is therefore reasonable to wonder whether the non-linearities induced by this non-linear mapping would somehow affect the conclusions presented here. A simple way to sidestep the present question is to undo the non-linear mapping by defining a new observable $\\hat{\\flux}=-\\ln(\\flux)=A(1+\\delta_ {\\rm IGM})^{\\beta}$ and proceed by measuring its correlations.  The best way to assess to what extent the above caveats affect the estimates reported in the present work is through numerical simulations, calculating the convergence field on a light cone and at the same time measuring \\Lya forest synthetic spectra and cross-correlating the two. This will be the next step in our investigation and the focus of the next publication.  Finally, on the analytical side we still need to address the estimate of the correlators when the power spectrum shows evolution in redshift \\textit{and} on different scales at the same time. As pointed out, $\\langle \\delta\\flux^{2} \\kappa\\rangle$ is sensitive to scales $k\\gtrsim 10^{-2} \\,\\hMpcI$. As such this correlator is an ideal tool to test modifications of gravity that show scale dependent growth. At the same time, this development would also allow the implementation of the hyperextended perturbation theory results and as such to address analytically the impact of gravity induced nonlinearities on the value of the correlators.   \\textit{Acknowledgements:} We thank S.~Matarrese, F.~Bernardeau, S.~Dodelson, J.~Frieman, E.~Sefusatti, N.~Gnedin, R.~Scoccimarro, S.~Ho, D.~Weinberg and J.~P.~Uzan for useful conversations. AV is supported by the DOE  at Fermilab. MV is supported by grants PD51, ASI-AAE and a PRIN MIUR. DNS and SD are supported by NSF grant AST/0707731 and NASA theory grant NNX08AH30G. DNS thanks the APC (Paris) for its hospitality in spring 2008 when this project was initiated. AV thanks IAP (Paris) for hospitality during different stages of this project.   \\onecolumngrid  \\appendix"
    },
    "0910/0910.0126_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\setcounter{equation}{0}  Galaxy clusters are the largest virialized structure formed by the gravitational collapse of small density perturbations. The spherical collapse model (SCM) is a simple and powerful tool for investigating the evolution of nonlinear structure and bound systems in the Universe \\cite{Gunn,Peebles}.  Since the formation of the large scale dark matter potential well of clusters depends only on gravitation, the property of dark energy (DE) and the cosmological parameters can be constrained from the growth of large scale structure and the abundances of rich clusters of galaxies. There have been numerous papers investigating this problem with the spherical collapse model \\cite{0401504,9604141,WS,0409481,0504465,0505308,0507257,07080868,09063349,09081333}.  In order to prohibit the spherical overdensity from collapsing to a singular point in a finite time, we need to use the virialization in the SCM. The common mistake in literatures is that one uses the correct value of the finite virial radius of the overdensity from the virial theorem but still uses the incorrect value of the background scale factor for an Einstein de Sitter (EdS) universe. However, this leads to an incorrect conclusion that the bound systems are formed when the nonlinear overdensity is about $178$. Equally, this leads to the incorrect critical linear threshold density $1.68$ \\cite{sky}. When we use the correct values of the virial radius and the scale factor at the virial epoch, these values are lowered to $147$ and $1.58$, respectively.  In the next section, we review and analyze the SCM model in a flat universe including the non-clustering DE with the constant equation of state $\\omde$. We obtain the exact solution for the evolution of the scale factor (a) in terms of the ratio of it to its value at the turnaround epoch ($x \\equiv \\fr{a}{a_{\\ta}}$). We also obtain the exact (when $\\omde = -\\fr{1}{3}$) and the approximate (when $\\omde \\neq -\\fr{1}{3}$) solutions for the ratio of the overdensity radius to its value at the turnaround epoch ($y \\equiv \\fr{R}{R_{\\ta}}$) for general cosmological parameters. Also the exact and approximate solutions of the overdensity at the turnaround epoch $\\zeta = \\fr{\\rho_{\\mc}}{\\rho_{\\m}} \\Bigl|_{z=z_{\\ta}}$ are obtained for general $\\omde$. The given approximate solutions have the less than $1.5$ \\% errors over the parameter space in the concordance model. In Sec. $3$, we use the virial theorem to obtain the correct values of $x$ and $y$ at the virial epoch. We compare the values of the virial epoch and the nonlinear overdensity for the different models. We conclude in Sec. $4$. ",
        "conclusions": "\\setcounter{equation}{0}  The exact solution for the evolution of the background scale factor in a flat universe with the dark energy is given as a function of time. This provides the exact value of collapsing time and makes it possible to test the accuracies of both the analytic and the numerical solutions. We obtain the exact and approximate solutions of the evolution of the overdensity radius with high accuracy. From these solutions we are able to estimate the non-linear overdensity at any epoch.  Even though the analytic forms of $y$ and $\\zeta$ are obtained for the constant $\\omde$ models, they can be generalized to the slowly varying $\\omde \\geq -\\fr{1}{3}$ because they provide the very accurate values of them for the wide range of cosmological parameters.  The virial theorem provides the exact value of the virial epoch for given models. From this, the value of the non-linear overdensity is obtained. In the high matter density and/or the high-redshift clusters, the model dependence of the observed quantities becomes weaker. We are able to categorize the physical quantities of the low and high redshift clusters for the given models.  From the correct value of the virial radius of an EdS universe, we obtain the correct linear and non-linear overdensities for an EdS universe, which are lower than the known values. In this paper, we show that clusters are formed earlier than those from the conventional model. Thus, the mass of the cluster is larger than one obtained from the conventional model. This will affect to the prediction for the different models for the cluster abundances, the weak gravitational lensing, etc.  \\appendix"
    },
    "0910/0910.0310_arXiv.txt": {
        "abstract": "The Aquarius project is the first simulation that can resolve the full mass range of potential globular cluster formation sites.  With a particle mass $m_\\mathrm{p}=1.4 \\times 10^4$\\Msun, Aquarius yields more than 100 million particles within the virial radius of the central halo which has a mass of $1.8 \\times 10^{12}$\\Msun, similar to that of the Milky Way. With this particle mass, dark matter concentrations (haloes) as small as 10$^6$ M$_\\odot$ will contain a minimum of 100 particles.  Here, we use this simulation to test a model of metal-poor globular cluster formation based on collapse physics. In our model, globular clusters form when the virial temperatures of haloes first exceed $10^4$\\,K as this is when electronic transitions allow the gas to cool efficiently.  We calculate the ionising flux from the stars in these first clusters and stop the formation of new clusters when all the baryonic gas of the galaxy is ionised. This is achieved by adopting reasonable values for the star formation efficiencies and escape fraction of ionising photons which result in similar numbers and masses of clusters to those found in the Milky Way. The model is successful in that it predicts ages (peak age $\\sim$ 13.3 Gyrs) and a spatial distribution of metal-poor globular clusters which are consistent with the observed populations in the Milky Way. The model also predicts that less than 5$\\%$ of globular clusters within a radius of 100 kpc have a surviving dark matter halo, but the more distant clusters are all found in dark matter concentrations.  We then test a scenario of metal-rich cluster formation by examining mergers that trigger star formation within central gas disks. This results in younger ($\\sim$ 7--13.3 Gyrs), more centrally-located clusters (40 metal-rich GCs within 18 kpc from the centre of the host) which are consistent with the Galactic metal-rich population. We test an alternate model in which metal-rich globular clusters form in dwarf galaxies that become stripped as they merge with the main halo. This process is inconsistent with observed metal-rich globulars in the Milky Way because it predicts spatial distributions that are far too extended. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.2854_arXiv.txt": {
        "abstract": "\\noun{fish} is a fast and simple ideal magneto-hydrodynamics code that scales to $\\sim 10\\,000$ processes for a Cartesian computational domain of $\\sim 1000^{3}$ cells. The simplicity of\\noun{fish} has been achieved by the rigorous application of the operator splitting technique, while second order accuracy is maintained by the symmetric ordering of the operators. Between directional sweeps, the three-dimensional data is rotated in memory so that the sweep is always performed in a cache-efficient way along the direction of contiguous memory. Hence, the code only requires a one-dimensional description of the conservation equations to be solved. This approach also enable an elegant novel parallelisation of the code that is based on persistent communications with MPI for cubic domain decomposition on machines with distributed memory. This scheme is then combined with an additional OpenMP parallelisation of different sweeps that can take advantage of clusters of shared memory. We document the detailed implementation of a second order TVD advection scheme based on flux reconstruction. The magnetic fields are evolved by a constrained transport scheme. We show that the subtraction of a simple estimate of the hydrostatic gradient from the total gradients can significantly reduce the dissipation of the advection scheme in simulations of gravitationally bound hydrostatic objects. Through its simplicity and efficiency, \\noun{fish} is as well-suited for hydrodynamics classes as for large-scale astrophysical simulations on high-performance computer clusters. In preparation for the release of a public version, we demonstrate the performance of \\noun{fish} in a suite of astrophysically orientated test cases. ",
        "introduction": "\\label{sec:introduction}  It has been argued that computational models provide the third pillar of scientific discovery, beside the traditional experiment and theory. On the practical side, there are indeed a lot of similarities between computer models and experiments: In both cases one tries to start from a well-defined initial state. The outcome of the simulation is in the same way unknown as the outcome of an experiment and unexpected results are interesting in both cases. Also, one has to define observables whose values one intends to measure at several time intervals or at the end of the simulation or experiment. If the computer model or experiment is complicated enough, these values will be affected by systematic and statistical errors and their understanding requires an involved theoretical analysis and interpretation. Finally, both computer model and experiment need sufficient, and sometimes quite extensive, documentation to make the results reproducible.  On the other hand, only the real experiment is capable of probing the properties of nature, while the computer model is bound to elaborate the properties of the theory it is built on. The strength of the computer simulation is that it may demonstrate and extend the predictive power of its underlying theory. In principle, though, it seems completely irrelevant whether the consequences of the underlying theory are evaluated by analytical approaches or a computer-aided model. The computer model can be regarded as the analytical evolution of a discretised model of the physics equations, where the computer only provides appreciated efficiency in the evaluations, but does not add anything new to an analytical investigation.  This leads to our point of view that the computer model is not an independent pillar of science, but an increasingly important contemporary method in the development of theory. The computer model is most helpful if many different physics aspects are intertwined in a manner that prevents the treatment of the problem by dividing it into subproblems. Classical examples are found in research areas whose subject is inaccessible to human manipulation and preparation in the laboratory, like meteorology, planetary interiors, stars, and astrophysical dynamics in general.  Because physics on many different length and time scales is involved, it is of primary importance to first perform order of magnitude estimates to identify irrelevant processes and to find equilibrium conditions that constrain degrees of freedom. A classical theoretical model is then built by the description of a sequence of dominant processes, supported by order of magnitude estimates. The next step in the traditional approach is to approximate the complicated microscopic input physics by simple analytical fitting formulas that depend on few parameters. If the fitting formulas are simple enough, the problem might become analytically solvable so that the dependence of the global features on the chosen parameters can be revealed. In stellar structure, for example, order of magnitude arguments show that the matter is in thermodynamical equilibrium. Hence, an equation of state can be defined that expresses the gas pressure, $p$, as a function of the local density, $\\rho$, temperature and composition. In some regimes, the microscopic physics information in the equation of state is well approximated by an equation of state of the form $p=\\kappa\\rho^{1+1/n}$. Here, $\\kappa$ is a constant and $n$ a parameter called the polytropic index. If this approximation holds, the equations of stellar structure can be integrated analytically for selected integer values of $n$ \\cite{1967aits.book.....C,Lane1869}. Hence, traditional theoretical astrophysics relies on two prongs: (1) the interpretation of observations in terms of scenarios that are supported by order of magnitude estimates, and (2) the explanation of complicated processes by simple approximate laws that allow a precise analytical investigation of the most important aspects of the model.  It is obvious how scientific computing can extend the lever of approach (2). As soon as the complicated processes are reduced to simple approximate laws, the initial state and the equations of evolution of the model are mathematically well-defined. It is therefore possible to replace the former analytical investigation by calculations on the computer. For example, the Lane-Emden equation can be integrated numerically for many more values of the polytropic index. There is only one mathematically correct answer to each problem and the numerical calculation can be refined until the algorithm obtains the desired precision of the result.  It is much more intricate to let computer models contribute to approach (1), i.e. the interpretation of observations in terms of scenarios that are supported by order of magnitude estimates. A well-known example is the propagation of a shock front. Only a minority of computer codes attempt to resolve the microscopic width of the shock to correctly treat its complicated input physics. Most codes just ensure the accurate conservation of mass, momentum and energy across the shock front to obtain the shock propagation speed and the thermodynamic conditions on both sides of the shock. Here, it is the numerical algorithm that dynamically performs the simplification of the model and not the description of the input physics like in approach (2). Hence, one may implement a rich set of input physics and design the finite differencing such that fundamental laws of physics are fulfilled under all possible conditions. Then, the complexity of the model is bound by the scale on which unresolvable small-scale structures are dissipated in space and time. Because this scale depends on the numerical algorithm and the mesh topology rather than the investigated physics, different solutions of the same physics problem may not converge to the same microscopic solution. However, the goal of approach (1) is not to evaluate the detailed microscopic state of $\\sim10^{9}$ individual fluid cells (for a human being it would take several years of uninterrupted reading to absorb this amount of information), but to elaborate the generic macroscopic properties that are determined by the fundamental physics laws in combination with the microscopic input physics. The actual microscopic state of the fluid cells acts as an exemplary sampling of a specific macroscopic state. Convergence of the evolution of the macroscopic state must be demonstrated by careful comparison and theoretical interpretation of concurrent implementations of the same physics problem using different algorithms and meshes. Hence, we believe that it is important for the reliability of astrophysical models to be investigated with several different radiation-magneto-hydrodynamics codes that are simple and efficient enough to be broadly used.  There are several well-documented and publicly available software packages. \\noun{gadget} is a cosmological N-body code based on the Smoothed Particle Hydrodynamics method and parallelised using the Message Passing Interface (MPI). It is a Lagrangean approach with a hierarchical tree for the non-local evaluation of Newtonian gravity \\cite{2001NewA....6...79S}. \\noun{vh-1} is a multidimensional hydrodynamics code based on the Lagrangian remap version of the Piecewise Parabolic Method (PPM) \\cite{1993ApJS...88..589B} that has been further developed and parallelised using MPI. \\noun{zeus-2D} \\cite{1992ApJS...80..819S,1993ApJ...413..198S,1993ApJ...413..210S} and \\noun{zeus-3D} are widely used grid-based hydrodynamics codes for which a MPI-parallel \\noun{zeus-mp} version exists as well. It offers the choice of different advection schemes on a fixed or moving orthogonal Eulerian mesh in a covariant description. While these more traditional approaches distribute in the form of a software package that includes options to switch on or off, recent open source projects try to provide the codes in the form of a generic framework that can host a variety of different modules implementing different techniques. In this category we could mention the \\noun{flash} code \\cite{2002ApJS..143..201C,2000ApJS..131..273F} with its main focus on the coupling of adaptive mesh stellar hydrodynamics to nuclear burning and the \\noun{whisky} code \\cite{2007CQGra..24..235G} as a recent general relativistic hydrodynamics code based on the \\noun{cactus} environment. Further we mention \\noun{athena} \\cite{0067-0049-178-1-137} and \\noun{enzo} \\cite{2004astro.ph..3044O} which are both well documented and publicly available codes.  In this paper we present our code \\noun{fish} (Fast and Simple Ideal magneto-Hydrodynamics), which follows a somewhat different strategy. \\noun{fish} is based on the publicly available serial version of a cosmological hydrodynamics code \\cite{2003ApJS..149..447P}. Its main virtue is the simplicity and the straightforward approach on a Cartesian mesh. It can equally well be used for educational purpose in hydrodynamics classes than for three-dimensional high-resolution astrophysical simulations \\cite{2003ApJ...596L.207P}. However, in the first code versions, simplicity and parallel efficiency were not available at the same time. Here we describe a new and elegant implementation of the parallelisation with a hybrid MPI/OpenMP approach that has been redesigned from scratch and is well separated from the subroutines describing the physical conservation equations. Hence, the same code version is now simple \\emph{and} efficient on large high-performance computing clusters. In comparison to the above-mentioned software packages it is simple to modify and/or extend and has negligible overhead. With respect to earlier versions, the discretisation of the advection terms and the implementation of gravity has further been developed to make the \\noun{fish} code more accurate in the treatment of fast cold flows and more stable for the evolution of gravitationally bound hydrostatic objects. Special attention was given to the robustness of the approach in regions with poor resolution, which are difficult to avoid in global astrophysical models. \\noun{fish} has successfully been used in one of the first predictions of the gravitational wave signal from 3D supernova models with microscopic input physics \\cite{2008A&A...490..231S} and provides the foundation for the implementation of spectral neutrino transport in our new code \\noun{elephant} (ELegant and Efficient Parallel Hydrodynamics with Approximate Neutrino Transport) \\cite{2009ApJ...698.1174L}.  In Section 2 we describe the physics equations and document the numerical methods used in \\noun{fish}. In Section 3 we discuss the optimisation of the memory access, the parallelisation of \\noun{fish} with MPI and OpenMP and present the scaling behaviour to $\\sim10000$ processes. In Section 4 we finally investigate the performance of \\noun{fish} in a suit of six mostly magneto-hydrodynamic test problems. ",
        "conclusions": "\\label{sec:conclusion} We describe a simple and efficient three-dimensional implementation of the ideal magneto-hydrodynamics equations in our code named \\noun{fish}. The algorithms used are mostly based on \\cite{2003ApJS..149..447P}. Three different choices make the code very transparent and easy to work with: First, the rigorous operator splitting in second order operator ordering separates the sweeps across different coordinate directions. Also the constrained transport of the magnetic field is operator split from the hydrodynamics update. Second, the use of a Riemann solver free relaxation scheme with second order TVD advection facilitates the update of the state vector. Second order in time is achieved by a predictor-corrector approach within each time step. Third, subroutines for the physics operators are only required for the $x$-direction because the sweeps in $y$-direction and $z$-direction are performed by a rotation of the computational domain. Between the sweeps, the data is rotated such that an $x$-sweep in the direction of contiguous memory actually performs a $y$- or $z$-sweep in the coordinate frame of the rotated data. We document the details of an improved discretisation of the advective fluxes with respect to earlier versions of the code. Due to its clarity, we have successfully used the code for exercises and computational experiments in magneto-hydrodynamics classes.  We present different tests to evaluate the accuracy of \\noun{fish}. Second order convergence is demonstrated by the example of a circularly polarised propagating Alfven wave. The capability of accurately capturing hydrodynamic shocks is shown in a Sedov-Taylor blast wave explosion and in a spherical Riemann problem that are evolved in the three-dimensional computational domain. The latter two examples produce nice spherical flow features and show that the directional operator splitting does not cause sizeable asymmetries. We also investigate strong magnetosonic shocks in the form of a magnetic explosion. For this problem we reached a limit of the current scheme for very low plasma $\\beta$ and very strong shocks. However, there is still room for improving the scheme in this extreme regimes and future developments will go in that direction. In order to further test the coupling between the magnetic fields and the fluid we perform the rotor problem in a two-dimensional computational domain. Finally, we calculate a 2D MHD Riemann problem and successfully compare the result with the literature. The only test that led to insufficient performance was the time evolution of a gravitationally bound hydrostatic object at low resolution. The problem is that the presence of gravitation leads to stationary density gradients that are mistaken as propagating wave gradients in the scheme that tailors the numerical dissipation for a total variation diminishing solution. The corresponding stationary dissipation leads to the evaporation of the bound object on few dynamical time scales. We remedied this problem by analytically estimating the hydrostatic gradients and subtracting them from the total gradients before calculating and applying the dissipation for the propagating waves. This to our knowledge novel and simple measure could also be used in combination with the Roe Riemann solver and potentially other approximate solvers. It allows the stable evolution of a gravitationally bound object over many dynamical time scales without exaggerated numerical dissipation at the center. This is shown in our last test case where we evolve a polytrope in hydrostatic equilibrium.  Inspite of the straightforward approach, \\noun{fish} is very efficient, even for problems containing $\\sim 1000^{3}$ and more cells. \\noun{fish} implements cubic domain decomposition for distributed memory computer clusters. The process-wise rotation of the data between the sweeps takes about $10\\%$ of the computational cost, but is largely compensated for by the fact that the individual sweeps along contiguous memory can be pipelined in the cache. The rotation of the data before the directional sweeps enables a very elegant parallelisation scheme that is based on non-blocking persistent MPI communication directives. We perform the communication for each direction separately and overlap it with computations of the update of interior cells. As sweeps along the same direction are embarassingly parallel within a shared memory cluster, we additionally parallelise those using OpenMP directives. With this, \\noun{fish} can take advantage of multi-core processors and scales nicely to $\\sim 10000$ processes for a problem of fixed size. The limiting factor in the scaling of \\noun{fish} is the increasing ratio of work on buffer zones to work on enclosed zones. Even more processes can be used efficiently if the problem size grows with the number of available processes.  \\noun{fish} is mainly targeted for astrophysical applications where simple and efficient low-order algorithms are appreciated to accept manifold variations of input physics. The input physics is often a patchwork of different descriptions and algorithms for the different domains. As they do not always connect with a prescribed smoothness and as decisions about approximations are often taken during runtime, it is difficult to work with higher order methods. \\noun{fish} has successfully been used in one of the first predictions of the graviational wave signal from 3D supernova models with microscopic input physics \\cite{2008A&A...490..231S} and provides the foundation for the implementation of spectral neutrino transport in our new code ELEPHANT (ELegant and Efficient Parallel Hydrodynamics with Approximate Neutrino Transport) \\cite{2009ApJ...698.1174L}. The dynamical contrast of the density in many astrophysical applications suggests the implementation of an adaptive mesh. We discussed the extension of the approach to non-uniform orthogonal meshes, but find them only applicable for limited dynamical contrasts of about a factor of two increase of the resolution. On the long term, it appears more promising to embed the uniform Cartesian mesh of our current approach in a multi-patch driver like e.g. carpet in the cactus framework. We hope that this paper will serve as a reference for future users and developers of the public version of \\noun{fish}, which is in preparation and will soon become available."
    },
    "0910/0910.0256_arXiv.txt": {
        "abstract": "In the low density intergalactic medium (IGM) that gives rise to the \\lyman\\ forest, gas temperature and density are tightly correlated.  The velocity scale of thermal broadening and the Hubble flow across the gas Jeans scale are of similar magnitude ($H\\lambda_J \\sim \\sigma_{\\rm th}$). To separate the effects of gas pressure support and thermal broadening on the \\lya\\ forest, we compare spectra extracted from two smoothed particle hydrodynamics (SPH) simulations evolved with different photoionization heating rates (and thus different Jeans scales) and from the pressureless dark matter distribution, imposing different temperature-density relations on the evolved particle distributions. The dark matter spectra are similar but not identical to those created from the full gas distributions, showing that thermal broadening sets the longitudinal (line-of-sight) scale of the \\lya\\ forest. The turnover scales in the flux power spectrum and flux autocorrelation function are determined mainly by thermal broadening rather than pressure. However, the insensitivity to pressure arises partly from a cancellation effect with a sloped temperature-density relation ($T \\propto \\rho^{0.6}$ in our simulations): the high density peaks in the colder, lower pressure simulation are less smoothed by pressure support than in the hotter simulation, and it is this higher density gas that experiences the strongest thermal broadening.  Changes in thermal broadening and pressure support have comparably important effects on the flux probability distribution (PDF), which responds directly to the gas overdensity distribution rather than the scale on which it is smooth. Tests on a lower resolution simulation ($2\\times 144^3$ vs.\\ $2\\times 288^3$ particles in a $12.5\\hmpc$ comoving box) show that our statistical results are converged even at this lower resolution. While thermal broadening generally dominates the longitudinal structure in the \\lya\\ forest, we show in Paper~II that pressure support determines the transverse coherence of the forest observed towards close quasar pairs. ",
        "introduction": "\\label{sec:intro} The \\lyman\\ forest, caused by \\lya\\ absorption of neutral hydrogen atoms along the line of sight to some distant source (usually a quasar), was originally described in terms of discrete intervening gas ``clouds'' \\citep{lynds71,sargent80}, analogous to clouds in the Galactic interstellar medium.  In this picture, the velocity widths of the observed absorption structures were primarily a consequence of thermal motions of the absorbing atoms.  In the mid-1990s, three-dimensional hydrodynamic simulations of cold dark matter (CDM) cosmological models achieved remarkable success in reproducing the observed properties of the \\lya\\ forest \\citep{cen94,zhang95,hernquist96,miralda96,theuns98}. In these simulations, most of the absorbing gas is at low density ($\\rho/\\bar{\\rho}\\sim0.1$--10) and naturally described as a continuously fluctuating medium rather than a series of discrete structures \\citep{hernquist96,bi97,rauch97,croft97,croft98}.  The velocity width of individual features is set largely by the Hubble flow across them, reflecting the physical extent of the absorbing gas along the line of sight \\citep{hernquist96,weinberg97}.  The temperature of the intergalactic medium (IGM), therefore, affects the structure of the \\lya\\ forest in two ways: by smoothing absorption along the line of sight through the thermal motions of atoms, and by smoothing the physical distribution of the gas in three dimensions through pressure support. In this study, we use two smoothed particle hydrodynamics (SPH) simulations with different photoionization heating rates---and thus different IGM temperatures and amounts of gas pressure support---to disentangle the relative effects of pressure support and thermal broadening in the \\lya\\ forest.  Both hydrodynamical simulations \\citep{katz96,miralda96,theuns98} and analytic arguments \\citep{hui97b} suggest that low-density intergalactic gas should have a power-law temperature-density relation, i.e., $T=T_0(1+\\delta)^{\\alpha}$, where $1+\\delta\\equiv\\rho/\\bar{\\rho}$ is the local gas overdensity.  Reasonable assumptions for heating and cooling rates yield $T_0\\sim 10^4$\\,K and $\\alpha\\sim0.6$ \\citep{hui97b,theuns98}.  Observations, however, imply that the normalization $T_0$ could be nearly twice as high, and the slope could be much shallower or even inverted \\citep{schaye99,ricotti00,mcdonald01,bolton08}.  Regardless of the parameter values, because the optical depth to \\lya\\ absorption is related to both the temperature and the density by a power law, the existence of the temperature-density relation (also called the ``equation of state'') implies a tight relation between the observed \\lya\\ absorption and the IGM gas density.  Since the universe has $\\sim 5$ times as much mass in dark matter as in baryons, we can expect in general for intergalactic gas to trace the underlying dark matter on scales above the gas Jeans length \\citep{schaye01}.  The \\lya\\ forest therefore provides a powerful tool for tracing the dark matter power spectrum in the quasi-linear regime---modulo the effects of peculiar velocities, thermal broadening, and gas pressure \\citep{croft98,croft99,croft02,mcdonald00,mcdonald06,viel04,viel06a}.  Pressure should be important on scales at or below the Jeans length, \\begin{equation}\\label{eqn:jeans} \\lambda_J = c_s\\sqrt{\\frac{\\pi}{G\\rho}} = \\sigma_{\\rm th}\\sqrt{\\frac{5\\pi}{3G\\rho}}, \\end{equation} where $c_s=\\sqrt{[5kT]/[3m]}=\\sigma_{\\rm th}\\sqrt{5/3}$ is the speed of sound in an ideal gas expressed as a multiple of the 1-D thermal velocity $\\sigma_{\\rm th}$ \\citep{miralda96,schaye01,desjacques05}.  Here $m$ is the average mass of the gas particle; in an ionized primordial mixture of hydrogen and helium, $m=0.59m_p$, where $m_p$ is the proton mass. The density $\\rho$ is the density of the gravitating medium, which is dominated by dark matter, so we take $\\rho = \\Omega_{m,0}\\rho_{c,0}(1+z)^3 (1+\\delta)$.  In comoving coordinates, \\begin{eqnarray}\\label{eqn:jeanscomove} \\lambda_{J,{\\rm comv}} & = & (1+z)\\sigma_{\\rm th}H_0^{-1}\\sqrt{\\frac{5\\pi}{3}}\\left[\\frac{3}{8\\pi}\\Omega_{m,0}(1+z)^3(1+\\delta)\\right]^{-1/2}\\nonumber\\\\ &= & 782\\,h^{-1}\\,\\mbox{kpc}\\\\ &\\times& \\left(\\frac{\\sigma_{\\rm th}}{11.8\\,\\rm{km}\\,\\rm{s}^{-1}}\\right)\\left[\\left(\\frac{\\Omega_{m,0}(1+\\delta)}{0.25\\times(1+0)}\\right)\\left(\\frac{1+z}{1+3}\\right)\\right]^{-1/2}, \\nonumber \\end{eqnarray} where we have normalized the thermal broadening velocity $\\sigma_{\\rm th}$ to correspond to a fiducial temperature of $10^4$\\,K.  By defining a ``Jeans velocity'' as $v_J\\equiv H\\lambda_J$ we can find the relative importance of the Jeans scale and $\\sigma_{\\rm th}$, \\begin{eqnarray}\\label{eqn:vjsigth} \\frac{v_J}{\\sigma_{\\rm th}} &=& \\frac{2\\pi\\sqrt{10}}{3}(1+\\delta)^{-1/2}\\left[\\frac{\\Omega_{m,0}(1+z)^3 + \\Omega_{\\Lambda}}{\\Omega_{m,0}(1+z)^3}\\right]^{1/2}\\\\\\nonumber &\\approx& 6.62(1+\\delta)^{-1/2}, \\end{eqnarray} which is essentially redshift-independent for the redshifts relevant to the \\lya\\ forest.  The Jeans length of equation~(\\ref{eqn:jeans}) divides stable from unstable modes in a static, homogeneous, self-gravitating medium.  The IGM is expanding, inhomogeneous, non--self-gravitating (because dark matter dominates), and evolving in density and temperature on the same timescale that fluctuations grow. Even for linear perturbations in a baryonic universe, the ``filtering scale'' below which fluctuations growth is suppressed depends on the thermal history of the gas rather than the instantaneous temperature-density relation \\citep{gnedin98}.  We therefore expect equation~(\\ref{eqn:jeans}) to describe the scale of gas pressure support only at an order-of-magnitude level.  Equation~(\\ref{eqn:vjsigth}) shows that Jeans velocities and thermal velocities should be comparable at the overdensities of typical \\lya\\ forest features, but the calculation is not definitive enough to show whether one will dominate in practice.  Hence, to predict the statistical properties of the \\lya\\ forest in a given cosmological model, one must calculate the predicted gas distribution.  The most reliable way to do this uses full $N$-body plus hydrodynamic simulations, which include the gravity of dark matter and gas and the additional effects of pressure, adiabatic heating and cooling, heating by photoionization and shocks, and radiative cooling.  However, large volume hydrodynamic simulations with the necessary resolution are computationally intensive. \\citet{weinberg97} and \\citet{croft98} show that one can achieve reasonable accuracy in the \\lya\\ forest regime from pure dark matter simulations, applying the temperature-density relation to the evolved dark matter distribution to compute spectra.  Indeed, the log-normal model \\citep{bi97}, in which the dark matter distribution is computed from the linear density field by an exponential transformation \\citep{coles91}, provides a qualitatively accurate physical model of the \\lya\\ forest, sufficient for creating artificial spectra with reasonable statistical properties.  A fast way to include the effects of gas pressure in an approximate way is the hydrodynamic particle-mesh (HPM) method \\citep{gnedin98,ricotti00,meiskin01}.  HPM assumes that all the gas follows the power-law IGM equation of state, and it uses a modification of the standard (fast) particle-mesh $N$-body method to compute the sum of gravitational forces and pressure gradient forces given this equation of state.  In a detailed comparison of fully hydrodynamic SPH simulations and the approximated HPM simulations, both evolved using {\\sc Gadget-2}, \\citet{viel04} found that while the HPM approach does converge to the SPH results, for some \\lya\\ forest properties, such as the flux probability distribution or small-scale power spectrum, it can differ from the SPH calculation by as much as 50\\% at $z\\sim 2$.  A simpler alternative to HPM is to use a pure $N$-body simulation but smooth the evolved dark matter distribution on the Jeans scale before extracting \\lya\\ forest spectra \\citep{zaldarriaga01,desjacques05}. Because the Jeans smoothing is three-dimensional, it is not degenerate with line-of-sight thermal broadening. \\citet{zaldarriaga01} found that, even though the thermal broadening dominates the pressure correction, the value of the Jeans length becomes a large source of uncertainty in cosmological inferences from the \\lya\\ forest if one tries to estimate it from the data rather than predict it from theory.  Because the Jeans length is expected to vary with the gas temperature and density, a Zel'dovich-like scheme can be used when smoothing the dark matter distribution in order to model these effects \\citep{viel02a}.  Though we do not carry out a comprehensive comparison of these methods here, we do investigate the impact of pressure in detail, as well as compare full SPH simulations to results derived from the dark matter distribution alone.  The primary purpose of this paper and its companion (Peeples et al.\\ 2009, hereafter Paper~II) is to disentangle the roles of pressure support and thermal broadening in the \\lya\\ forest by studying two SPH simulations with different thermal histories and temperature-density relations.  In addition to examining the physics of the \\lya\\ forest, this study is motivated by observational evidence (discussed in \\S\\,\\ref{sec:rhot}) that the temperature of the IGM at $z\\sim 2$--4 is higher than expected from simple photoionization models by a factor of 1.5--2.  This evidence comes from analyses of data along single lines of sight.  Because the higher pressure associated with hotter gas would smooth the IGM in three dimensions, using closely paired lines of sight to probe this coherence scale has been proposed as an alternative route for inferring the temperature-density relation (J.\\ Hennawi, private communication, 2007).  However, before we can understand the relative roles of temperature and pressure on the transverse structure of the \\lyman\\ forest, we must first understand their longitudinal effects along independent sightlines, which is the goal of this paper.  While here we find that thermal broadening dominates pressure support in setting the level of longitudinal structure in the \\lya\\ forest, in Paper II we show that the gas Jeans length dominates the level of coherence transverse to the line of sight.  This paper is organized as follows.  In \\S\\,\\ref{sec:sims}, we describe the SPH simulations we have evolved to investigate these effects.  In \\S\\,\\ref{sec:physics}, we describe the physics of the \\lya\\ forest in these simulations and examine the impact of the thermal history on observable spectra; this section also serves to review the physical understanding of the high-redshift \\lya\\ forest that has emerged from simulations and associated analytic work since the mid-1990s. We then study these effects on several typical statistical measures for learning about the IGM from the \\lya\\ forest in \\S\\,\\ref{sec:stats}, with our conclusions in \\S\\,\\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc} We have investigated the relative importance of pressure support and thermal broadening in determining the longitudinal structure of the \\lyman\\ forest.  Our main results come from comparing two SPH simulations with identical initial conditions but different photoionization heating rates, which produce a factor of $\\sim 2.3$ difference in the temperature of diffuse IGM gas.  We have imposed different temperature-density relations on the simulation outputs to isolate physical effects, extracted spectra from the dark matter distribution to extend our investigation to the pressureless case, and compared the fiducial simulation ($2\\times 288^3$ particles) to a lower resolution simulation ($2\\times 144^3$) to quantify numerical resolution effects.  Equation~(\\ref{eqn:vjsigth}) shows that the Hubble flow across the Jeans scale is generally of the same magnitude as thermal broadening ($H\\lambda_J\\sim\\sigma_{\\rm th}$).  However, the IGM is an expanding, inhomogeneous medium evolving in the potential of a non-linear dark matter distribution, so the Jeans length is at best an approximate description of the scale imposed by gas pressure support.  It is therefore difficult to know without detailed simulations whether thermal broadening or pressure support will dominate the structure of the forest.  A related question is the meaning of the density contrast $\\delta$ in the fluctuating Gunn-Peterson approximation (equation~\\ref{eqn:tau}).  In principle, this should be the density of the gas that is absorbing \\lya\\ photons, but analyses of early SPH simulations found that \\lya\\ forest spectra created from the (pressureless) dark matter distribution were remarkably similar to those created from the (pressure supported) gas distribution \\citep[e.g.,][]{croft98}. However, the relatively low resolution of those simulations (a mass resolution that is $\\sim 60$ times lower than our fiducial simulation here) raised the possibility that both sets of spectra were artificially broadened to the numerical resolution limit. We investigate this issue more confidently here because the initial particle spacing of our $288^3$ simulations, $43\\,h^{-1}$\\,kpc comoving, is far below the typical Jeans length at IGM overdensities, $\\lambda_J\\sim\\, 800h^{-1}$\\,kpc comoving, and because our $144^3$ simulation allows a direct resolution test.  In broad brush, our conclusions are that thermal broadening dominates over pressure support in determining the visual appearance and statistical properties of the \\lya\\ forest, but that differences in gas pressure do have a noticeable effect.  The widths of absorption features are typically  set by Hubble flow across the absorbing structure, though thermal broadening does smooth out small scale corrugations to create coherent features (see Fig.~\\ref{fig:zoo}).  Once we include thermal broadening, the spectra created from dark matter distributions are similar to those created from the gas, even though the effects of pressure support on the gas density field are readily discernible.  Thus, one can drastically change the Jeans scale without drastically changing the forest.  However, dark matter spectra are visually and statistically distinguishable from gas spectra, more so than in the earlier generation of lower resolution SPH simulations.  Turning to individual statistics, we find that thermal broadening sets the turnover scale of the one-dimensional flux power spectrum, with the fiducial gas, H4 gas, and dark matter distributions producing similar power spectra if one imposes the fiducial or H4 \\rhot\\ relation on all three.  Similar conclusions hold for the coherence scale of the flux decrement autocorrelation function.  However, in both cases, the weak impact of pressure support partly owes to a cancellation effect that arises with a sloped \\rhot\\ relation: a higher pressure distribution has more ``Jeans broadening,'' but it has less thermal broadening because there is less high overdensity, high temperature gas.  When we impose a constant IGM temperature of $T=2\\times 10^4$\\,K, the differences among the three cases are more noticeable, though they are still small compared to the effects of thermal broadening.  For the flux decrement probability distribution function, thermal broadening and pressure support have effects of comparable magnitude, though thermal broadening is still somewhat more important.  Our $144^3$ and $288^3$ simulations yield similar results for all our statistics.  The resolution effects are larger at $z=4$ than at lower redshifts, where they have a small but noticeable impact on all the statistics.  For most purposes, the resolution of our $144^3$ simulations (in a $12.5\\,h^{-1}$\\,Mpc~comoving volume) is adequate.  Though it is a stronger player than pressure support in setting the scale of the longitudinal \\lya\\ forest, thermal broadening is an inherently one-dimensional phenomenon.  Because pressure acts in three dimensions, we expect it to play the main role in setting the {\\em transverse} coherence of the \\lya\\ forest across neighboring sightlines. Growing samples of binary quasars with separations of $\\Delta\\theta\\lesssim 10$\\arcsec\\ now make it possible to probe the expected Jeans scale \\citep{hennawi06,hennawi09}.  We show in Paper~II that the degree of transverse coherence on these scales is indeed sensitive to gas pressure support and insensitive to thermal broadening.  Observational studies of close quasar pairs can directly probe the scale of pressure support on the \\lyman\\ forest, providing new insights into the physical state and thermal history of the high-redshift intergalactic medium."
    },
    "0910/0910.4405.txt": {
        "abstract": "In this paper, we present a foreground analysis of the \\WMAP\\ 5-year data using the {\\fastica} algorithm, improving on the treatment of the \\WMAP\\ 3-year data in \\citet{bottino_etal_2008}.  We determine coupling coefficients between the \\WMAP\\ data and templates commonly used to trace the dominant foreground emissions, and use them to study the spectral properties of the diffuse emissions and subsequently to clean the data. We confirm the dependence of the values of these scaling factors on the extension of the mask used in the analysis and we again demonstrate some anomalies when the $Kp0$ mask is adopted that remain unexplained.  We also revisit the nature of the free-free spectrum with the emphasis on attempting to confirm or otherwise the spectral feature claimed in \\citet{dobler_etal_2008b} and explained in terms of spinning dust emission in the warm ionised medium. With the application of different Galactic cuts, the index is always flatter than the canonical value of $2.14$ except for the $Kp0$ mask which is steeper. Irrespective of this, we can not confirm the presence of any feature in the free-free spectrum.  We experiment with a more extensive approach to the cleaning of the data, introduced in connection with the iterative application of {\\fastica}. We confirm the presence of a residual foreground whose spatial distribution is concentrated along the Galactic plane, with pronounced emission near the Galactic center. This is consistent with the \\WMAP\\ haze detected in \\citet{finkbeiner_2004}.  Finally, we attempted to perform the same analysis on full-sky maps. The code returns good results even for those regions where the cross-talk among the components is high.  However, slightly better results in terms of the possibility of reconstructing a full-sky CMB map, are achieved with a simultaneous analysis of both the five \\WMAP\\ maps and foreground templates.  Nonetheless, some residuals are still present and detected in terms of an excess in the CMB power spectrum, on small angular scales. Therefore, a minimal mask for the brightest regions of the plane is necessary, and has been defined. ",
        "introduction": "\\label{intro} The most recent release of five years of observations from the \\WMAP\\footnote{\\emph{Wilkinson Microwave Anisotropy Probe}} satellite again allows a quantitative improvement in studies of the Cosmic Microwave Background (CMB) for cosmological purposes. However, such an improvement requires an analogous refinement in both our understanding of local astrophysical foregrounds and in the methods employed for the separation of these components from the CMB emission.  The most evident source of foreground contamination in the \\WMAP\\ data on at least large-angular scales is associated with the radiation produced by our own Galaxy -- due to synchrotron emission, free-free emission (or \\emph{thermal bremsstrahlung}) and thermal dust emission. The presence of an additional component is now also almost universally accepted -- this emission is strongly correlated with dust emission at microwave wavelengths but with a spectral dependence that is inconsistent with the basic thermal mechanism. It is generally referred to as \\emph{anomalous dust emission} since its nature is still not unambiguously identified.  Finally, a subject of ongoing debate is the so-called \\lq \\WMAP\\ haze' which \\citet{dobler_etal_2008a} attribute to \\emph{hard} synchrotron emission distributed around the Galactic Center, again with uncertain physical origin.  Many different techniques have been developed in order to clean the CMB data and subsequently study the foreground emission components. An inexhaustive selection of these includes: linear combination methods \\citep{bennett_etal_2003,tegmark_etal_2003,park_etal_2007,kim_etal_2008,delabrouille_etal_2008}, Gibbs sampling \\citep{eriksen_etal_2008},  the Maximum Entropy Method (MEM) \\citep{bennett_etal_2003,hinshaw_etal_2007}, WIFIT \\citep{hansen_etal_2006}, Correlated Component Analysis (CCA) \\citep{bonaldi_etal_2007}, and SMICA \\citep{delabrouille_etal_2003,patanchon_etal_2005,cardoso_etal_2008}. In addition, an implementation of the \\emph{Independent Component Analysis} approach \\citep{hyvarinen_1999,hyvarinen_oja_2000} referred to as \\fastica\\ has been successfully applied to the problem in \\citet{maino_etal_2002} and \\citet{baccigalupi_etal_2004}.  The aim of the current work is to undertake the foreground analysis of the \\WMAP\\ 5-year data with this algorithm, as previously applied to the \\WMAP\\ 3-year data \\citep{bottino_etal_2008}. As before, we focus on the diffuse Galactic foreground components and their spectral and spatial properties.  Specifically, we first apply \\fastica\\ to a set of data comprising one of the \\WMAP\\ frequency maps together with three standard templates of the diffuse Galactic emission, as described in section~\\ref{data} and section~\\ref{data_analysis}. We derive coupling coefficients for each frequency that, when used to scale the template amplitudes, are interpreted as the amount of foreground contamination for each foreground component at that frequency.  Then, we iteratively apply the algorithm to the set of data, cleaned according to these scale factors.  This second step of the analysis allows us to identify physical residuals that are not uniquely identified with one of the three templates.  Novel aspects of the analysis introduced in this paper include:   \\begin{itemize}  \\item consideration of the new KQ85 and KQ75 masks introduced by the \\WMAP\\ science team with the 5-year data release  \\item an investigation into the connection between the mask applied to derive the coupling coefficients and  the residuals resulting from the iterative multi-frequency analysis when these coefficients are applied, itself as a function of applied mask  \\item an attempt to understand the physical nature of the residuals revealed by the iterative analysis.  \\item the definition of a minimal mask that supports accurate foreground removal  \\end{itemize}  A general summary of the paper is as follows.  After a brief review in section~\\ref{ica_templates} on how we use {\\fastica} to perform the analysis, in section~\\ref{data} we describe the data and templates adopted, pointing out the main differences with respect to previous work. Then, in section~\\ref{simulations} we calibrate the performance of the method using Monte Carlo simulations and by comparison to a simpler $\\chi^2$-based method.  In section \\ref{data_analysis}, we present the results of the template fits, and study the spectral behaviour of the foregrounds with particular emphasis on the free-free emission.  The scaled templates are then used to clean the data, the resultant properties of which are investigated in section \\ref{cleaned_data}. In subsequent sections the iterative application of {\\fastica} then attempts to improve the removal of foreground contamination from the data.  Finally, we experiment with full-sky analysis, which benefits from the simultaneous analysis of multi-frequency data and templates. Section~\\ref{discussion} summarises our main conclusions. ",
        "conclusions": "\\label{discussion}  In this paper, we have undertaken a foreground analysis of the \\WMAP\\ 5-year data using the  {\\fastica} algorithm, as previously applied to the \\WMAP\\ 3-year data in \\citet{bottino_etal_2008}. Various improvements in the implementation have allowed us to address several open questions from our previous work, and allowed some experimentation with the technique.  We used the code to perform a foreground fit of the data on a frequency-by-frequency basis, where the Galactic components are described by all-sky templates obtained at wavelengths where the corresponding emission mechanisms dominate. Specifically, we adopted the \\citet{finkbeiner_2003} H$\\alpha$-map as a template for the free-free emission, the \\citet{finkbeiner_etal_1999} FDS8 model for thermal dust emission, and the 408~MHz radio continuum all-sky map of \\citet{haslam_etal_1982} as utilized in the first year \\WMAP\\ analysis for the synchrotron emission. In the latter case, we have also considered the differenced of the \\WMAP\\ K- and Ka-band data, as preferred in their 3-year analysis.  The first step of the analysis is the computation of the coupling coefficients, which give an estimation of the contamination of the foreground emission level in the data. We know already that these coefficients depend on the extension of the mask used to exclude the most contaminated regions of the Galactic plane. We further investigated this dependence, thanks to the new $KQ85$ and $KQ75$ masks provided by the \\WMAP\\ science team. We confirmed the result which reinforces the idea of a spatial variation of the foreground emissions. Moreover, it suggests the component separation is driven by some key regions, rather than the extension of the mask itself. On the other hand, we reestablish some anomalies obtained when the $Kp0$ mask is adopted that remain unexplained.  We have considered the spectral behaviour of the derived scaling factors when the \\citet{haslam_etal_1982} data is used as the synchrotron template, since an interpretation of the analogous results derived using the K-Ka template is compromised by the fact that it is itself a mixture of foreground components.  We evaluated the spectral index for the synchrotron emission, the anomalous dust-correlated component, and the free-free emission.  In the first two cases, we found steeper, though statistically consistent, spectral behaviour as compared to previous analysis, e.g. \\citet{davies_etal_2006}.  We then focused our attention on the estimation of the free-free spectral index. A non-trivial dependence of the spectral behaviour on the extension of the mask confirms what was already found with the \\WMAP\\ 3-year data. However, an understanding of the behaviour is probably associated with two factors -- the spatial variation of the physical properties of the component (e.g. electron temperature which directly relates to the scale-factor associated with the H$\\alpha$ template) which  is mostly linked to specific structures close the Galactic plane, and the connection of the statistical properties of the foreground with the response of the algorithm as a function of the non-linear function used. The spectral behaviour is generally flat and even increasing with frequency, if the $p$ function is used together with the $KQ85$ and $KQ75$ masks. We have interpreted these spectral features as a consequence of a self-similarity of the Gaussian/non-Gaussian mixtures at each frequency. When the $g$ function is adopted, a more uniform behaviour, consistent with theoretical expectations, is obtained. Finally, in both the cases, the $Kp0$ mask gives a steeper spectral index than the expected value of $2.14$. We note that none of our results are compatible with the existence of a bump in the spectrum, as claimed  by \\citet{dobler_etal_2008b} and subsequently explained by the existence of spinning dust in the WIM that is traced morphologically by the H$\\alpha$ template.  A more extensive approach to the cleaning of the data has been introduced in connection with the iterative application of {\\fastica}. A set of coefficients obtained with a specific mask is used to remove the foreground contaminations in the maps, which are then internally analysed with \\emph{all} the Galactic cuts available. This new approach allowed us to check the capability of the iterative step to recover from a poor initial cleaning of the data and still yield a credible CMB estimate. A useful figure of merit for the quality of the results is then the corresponding CMB power spectrum as compared to the \\WMAP\\ estimation. In general, the iterative approach is very robust, although some not completely satisfactory results provide a evidence of limitations of the algorithm. Specifically, it seems that the algorithm provides poorer estimates of the CMB if applied on smaller regions of the sky.  Associated with the iterative analysis, we were able to determine the presence of a residual foreground whose spatial distribution is concentrated along the Galactic plane, with pronounced emission near the Galactic center. The extension of this residual is independent of the mask used, meaning that it is not just the effect of a poor component separation. However, it decreases in amplitude if the K and Ka band are not included in the input data set, and more so if the K-Ka map is used to compute the synchrotron coefficients. This indicates that the component has a falling spectrum with frequency, but whether it is a new physical component or simply a manifestation of a region with different spectral behaviour from the average is difficult to determine. In any case, it does confirm the utility of the K-Ka data as a better template to describe the low-frequency foreground mixture. In fact, this emission was already observed in the SMICA analysis of \\citet{patanchon_etal_2005} and is clearly consistent with the \\WMAP\\ haze of \\citet{finkbeiner_2004}. The putative model of its spatial distribution proposed by \\citet{dobler_etal_2008a} is found to be in reasonable agreement with the data, although some refinements are required.  Finally, we attempted a full-sky analysis of the same data set. Even though the code then challenged to work in regions of the sky where the cross-talk among the components is very high, we find that {\\fastica} is still able to achieve good results, particularly based on the iterative analysis of the cleaned data as originally implemented in \\citet{bottino_etal_2008}. We focused our attention on the CMB reconstruction, comparing the results with those produced by a simultaneous analysis of the multi-frequency \\WMAP\\ data with the templates. In this case, the result is slightly better. On the other hand, a direct comparison of the power spectra of the ICA CMB map with variants of the ILC approach proposed recently (HILC, NILC) suggests the need to include more spatial dependence in our analysis. Indeed, the excess of power observed on small angular scales is likely the signature of some residual structures along the Galactic plane. Consequently, to partially compensate for such structures, we defined a minimal mask to be used in a practical analysis of the derived CMB map. The resulting \\fastica\\ power spectra for our preferred data sets then agree remarkably with the best estimate of the CMB spectrum provided by the \\WMAP\\ team."
    },
    "0910/0910.4643_arXiv.txt": {
        "abstract": "We present optical {\\it ESO} time-series and UV archival ({\\it FUSE}, {\\it HST}, {\\it IUE}) spectroscopy of the H-rich central star of He 2-138. Our study targets the central star wind in a very young planetary nebula, and explores physical conditions that may provide clues to the nature of the preceding post-AGB super-wind phases of the star. We provide evidence for a dense, slowly accelerating outflow that is variable on time-scales of hours. Line-synthesis modelling (SEI and CMFGEN) of low and high ionization UV and optical lines is interpreted in terms of an asymmetric, two-component outflow, where high-speed high-ionization gas forms mostly in the polar region. Slower, low ionization material is then confined primarily to a cooler equatorial component of the outflow. A dichotomy is also evident at photospheric levels. We also document temporal changes in the weak photospheric lines of He 2-138, with tentative evidence for a 0.36-day modulation in blue-to-red migrating features in the absorption lines. These structures may betray 'wave-leakage' of prograde non-radial pulsations of the central star. These multi-waveband results on the aspherical outflow of He 2-138 are discussed in the context of current interest in understanding the origin of axi- and point-symmetric planetary nebulae. ",
        "introduction": "There are several motivations for studying the time-variable characteristics of PN central stars (CSPN): Chandra and XMM-Newton observations are providing new perspectives on the presence of hot bubbles in planetary nebulae formed from the interaction between central star fast winds and the ambient asymptotic giant branch (AGB) phase gas (e.g. Kastner et al. 2008). The origin of the X-ray emission remains uncertain, as does the consequence of a potentially spatially clumped central star fast wind. We have previously reported on far-UV {\\it FUSE} satellite spectroscopic time-series observations of NGC~6543 (Prinja et al. 2007); there we demonstrated in particular the incidence of coherent, systematically evolving structures in the fast wind of the H-rich central star. The empirical characteristics derived by Prinja et al. (2007) are clearly comparable to the signatures of large-scale structure in clumped radiation pressure-driven winds of OB stars. The post-AGB mass-loss via a fast wind is important not only because of its dynamical interaction with the nebula, but also due to its effect on the evolution of the central star.  An improved understanding of the variable nature of CSPN spectral line profiles is also important in studying binarity in these objects, and thus post-common envelope evolution. In this case recorded photometric and spectroscopic signatures may be complicated by the influence of, for example, variable photospheric structure and an inhomogeneous inner stellar wind region (see e.g. De Marco et al. 2008). Furthermore, complex changes in diagnostic spectral lines can impact on non-LTE model atmosphere analyses aimed at deriving fundamental parameters of CSPN, including revisions to mass-loss rates due to strongly clumped outflows (e.g. Kudritzki, Urbaneja {\\&} Puls 2006; Urbaneja et al. 2008).  Though the UV resonance lines are undoubtedly the most appropriate for probing CSPN fast winds (e.g. Perinotto 1987), there is unfortunately a dearth of suitable UV {\\it time-series} datasets ({\\it IUE}, {\\it FUSE}, or {\\it HST}) currently available for these objects. We explore in this paper therefore the alternative of exploiting optical time-series data that offer diagnostics of variable conditions in, both, the fast wind and close to the stellar surface. We focus in this respect on high-quality {\\it ESO} optical data of the young H-rich central star He~2$-$138.  \\subsection{He~2-138} This paper reports on optical time-series and archival UV data of the central star of He~2-138 (PK~320-09 1). The central star provides an interesting balance in that (i) its fast wind reveals intermediate strength, unsaturated UV resonance lines, (ii) the weaker optical photospheric lines are very symmetric and likely not 'contaminated' by wind effects, and (iii)  the wind densities at low velocities are high enough to provide some optical signatures of the outflow. The nebula itself has a complex knotted appearance, with bubble-like structures arranged in fairly symmetrical form, providing an overall elliptical morphology (see e.g. the HST WFPC2 H$\\alpha$ imaging survey of Sahai {\\&} Trauger 1998).  There is no recent published model atmosphere analysis of the central star, though a non-LTE study deriving a few basic parameters is given by Mendez et al. (1988); Adopted parameters are listed in Table~1 and include values derived from a new non-LTE analyses carried out during our present study of He 2-138.   \\begin{table} \\centering \\caption{He~2-138 central star parameters.} \\begin{tabular}{lll} \\hline Parameter & Value & Reference  \\\\ \\hline  Sp. type & Of (H-rich) & Mendez et al. (1998) \\\\ T$_{\\rm eff}$ & 29000 $\\pm$ 1000 K  & This study. \\\\ Log $g$ & 2.95 $\\pm$ 0.15  & This study \\\\ Log (L/L$_\\odot$) & 3.90 & This study \\\\ Distance & 3.5 kpc & Zhang (1995) \\\\ Radial velocity & $-$49 km s$^{-1}$ & Schneider et al. (1983) \\\\ \\hline  \\end{tabular} \\end{table}   Hutton {\\&} Mendez (1993) have documented visual magnitude changes in He 2-138 with amplitudes of between $\\sim$ 0.1 and 0.15 mag. The fluctuations occur over timescales of a few hours and are similar to those evident in IC~418 (e.g. Mendez et al. 1986). The magnitude changes are not obviously modulated or periodic and may relate to variable conditions in the near-photospheric (inner-wind) regions. In their search for radial velocity variations in Southern Hemisphere CSPN, Afar {\\&} Bond (2005) quote a greater than 99{\\%} probability that He 2-138 is a radial velocity variable. These changes occur over $\\sim$ daily time-scales, but without any convincing evidence for a period.  Our aim in this study is to determine optical and UV mass-loss parameters for an object that is in the early stages of the planetary nebula phase. We explore here the evidence for a structured outflow in He 2-138 and examine corresponding or unrelated signatures for photospheric variability. ",
        "conclusions": ""
    },
    "0910/0910.1299_arXiv.txt": {
        "abstract": "To study the possible origin of the huge helium enrichment attributed to the stars on the blue main sequence of $\\omega$\\,Centauri, we make use of a chemical evolution model that has proven able to reproduce other major observed properties of the cluster, namely, its stellar metallicity distribution function, age-metallicity relation and trends of several abundance ratios with metallicity. In this framework, the key condition to satisfy all the available observational constraints is that a galactic-scale outflow develops in a much more massive parent system, as a consequence of multiple supernova explosions in a shallow potential well. This galactic wind must carry out preferentially the metals produced by explosive nucleosynthesis in supernovae, whereas elements restored to the interstellar medium through low-energy stellar winds by both asymptotic giant branch (AGB) stars and massive stars must be mostly retained. Assuming that helium is ejected through slow winds by both AGB stars and fast rotating massive stars (FRMSs), the interstellar medium of $\\omega$\\,Centauri's parent galaxy gets naturally enriched in helium in the course of its evolution. ",
        "introduction": "\\label{sec:int}  Once referred to as `simple stellar populations', globular clusters (GCs) are not simple at all, in that they display significant and peculiar star-to-star abundance variations (see Gratton, Sneden \\& Carretta 2004 for a review). These are most likely the records of a complex star formation history. Indeed, the C-N, Na-O, Al-O and Al-Mg anticorrelations seen among both evolved and unevolved stars of individual clusters cannot be reconciled with a simple deep mixing scenario (Gratton et al. 2001; Cohen, Briley \\& Stetson 2002). Rather, they are naturally explained if one or more successive stellar generations form out of the ejecta from first generation stars that have undergone nuclear processing through proton capture reactions at high temperatures (Sneden et al. 2004; Carretta et al. 2006, and references therein). This evolutive picture is supported by the presence of multiple main sequences (MSs) in two massive Galactic GCs, $\\omega$\\,Cen and NGC\\,2808 (Bedin et al. 2004; Piotto et al. 2007). The bluer MSs are -- surprisingly -- more metal-rich than the red ones, or have the same iron content. This, at present, can be understood only in terms of an extreme helium enhancement in the blue population (Norris 2004; D'Antona et al. 2005; Piotto et al. 2005). Further evidence in favour of the existence of very helium-rich subpopulations in massive clusters comes from the peculiar morphology of the horizontal branches in some of them (Busso et al. 2007; Caloi \\& D'Antona 2007; Yoon et al. 2008; but see also Catelan et al. 2009). The excess in the far-UV flux detected for most of the massive clusters observed in M\\,87 by Kaviraj et al. (2007) can also be interpreted as a signature of extreme helium enrichment and would demonstrate that this phenomenon is not limited to our Galaxy.  It is common wisdom that key to understanding the origin of extreme-helium populations is the identification of extreme helium polluters. Though massive (approximately 5--10 M$_\\odot$) asymptotic giant branch (AGB) stars have been suggested as likely `culprits' for the chemical `anomalies' observed in GC stars since the work by Cottrell \\& Da Costa (1981), only recently detailed stellar modelling has provided yields compatible with the observations over a significant range of initial stellar masses (Pumo, D'Antona \\& Ventura 2008; Ventura \\& D'Antona 2008a,b). In Ventura \\& D'Antona's models, convection is modelled efficiently and a very fast AGB evolution is obtained, which results in a relatively low number of thermal pulses (TPs) and third dredge-up (TDU) episodes. This leads to a lower production of carbon and nitrogen and no appreciable increase of the overall CNO abundance in the envelope, an important point in view of the fact that in $\\omega$\\,Cen and other globulars the [C+N+O/Fe]\\footnote{Unless otherwise stated, chemical abundances throughout this paper are by number, except for $Y$ and $Z$, that indicate the mass fraction of helium and total metals, respectively. As usual, log\\,$\\varepsilon$(X)~$\\equiv$ 12\\,+\\,log\\,(X/H), while square brackets indicate logarithmic ratios relative to solar, [A/B]~$\\equiv$ log\\,(A/B)\\,$-$\\,log\\,(A/B)$_{\\odot}$.} ratio is constant within a factor of about 2 (Norris \\& Da Costa 1995; Smith et al. 1996; Ivans et al. 1999; Carretta et al. 2005; Cohen \\& Mel\\'endez 2005; but see Yong et al. 2009).  Up to now, chemical evolution models aimed at explaining the large helium abundance implied by the blue MS of $\\omega$\\,Cen with the ejecta of AGB stars have predicted an enormous increase in the total C+N+O content of the cluster, contrary to observations (e.g. Karakas et al. 2006). This failure has spurred the quest for alternative solutions. Decressin et al. (2007a,b) have proposed a scenario where the H-processed material lost by first generation fast rotating massive stars (FRMSs) through slow mechanical equatorial winds is retained in the cluster potential well, where it enters the formation of second generation stars. Matter expelled through fast polar winds and supernova (SN) explosions leaves the cluster instead. A major drawback with this scenario is that it cannot estimate from first principles the efficiency of meridional circulation to mix helium (and other chemicals) into the stellar envelope, nor the rate of mass loss through the outflowing disc (Renzini 2008). Furthermore, the helium-rich stars would form in very `hostile' surroundings, their birth being shortly followed by -- or even concomitant to -- multiple SN explosions.  To reproduce the correct ratio of first generation-to-second generation stars in both the massive star pollution scenario and the AGB pollution scenario, either a highly anomalous initial mass function (IMF) for first generation stars or a strong evaporation of low-mass stars with `normal' chemical enrichment have to be assumed (Decressin et al. 2007b; D'Ercole et al. 2008). While, in general, the first condition is difficult to justify both theoretically and observationally, in the case of $\\omega$\\,Cen there are fairly clear indications that the second generation of stars could have formed from the ejecta of surrounding field stellar populations in an initially much more massive system. The kinematical, dynamical and chemical properties of this cluster, in fact, are better understood if it is the surviving nucleus of an ancient dwarf galaxy captured and disrupted by the gravitational field of the Galaxy many Gyr ago (see, for instance, Dinescu, Girard \\& van Altena 1999; Gnedin et al. 2002; Bekki \\& Norris 2006; Romano et al. 2007; Bellazzini et al. 2008). Recent {\\it N}-body simulations of the dynamical evolution of two-population clusters show that a rapid loss of first generation stars should be expected early on in the cluster evolution as a consequence of cluster expansion in response to the dynamical heating from SN explosions (D'Ercole et al. 2008; see also Decressin, Baumgardt \\& Kroupa 2008). Early phases of violent relaxation also sensibly reduce the initial cluster's mass (Meylan \\& Heggie 1997, and references therein). $\\omega$\\,Cen, however, must have followed a different evolutionary path, with its putative progenitor galaxy sheding stars in trails while its orbit was degrading by approaching the Milky Way plane (Meza et al. 2005). The debris of this disruption process has been possibly identified through the kinematical feature imprinted in a sample of local, metal-poor halo stars (Dinescu 2002; Mizutani, Chiba \\& Sakamoto 2003).  In this paper, we deal with the chemical evolution of the system that once was/contained $\\omega$\\,Cen. In particular, in Section~\\ref{sec:omega} we propose an explanation for the origin of its helium-rich population, in the context of a model that has already proven able to reproduce the majority of the observational constraints available for its whole, complex stellar population. For sake of completeness, in Section~\\ref{sec:omega} we also examine, in our framework, the two alternative scenarios of He being overenriched either only by AGBs or by FRMSs. Our conclusions are presented in Section~\\ref{sec:end}, together with a discussion of the results. ",
        "conclusions": "\\label{sec:end}  It is common wisdom that $\\omega$\\,Cen is the naked nucleus of a dwarf spheroidal galaxy that was ingested and partly disrupted by the Milky Way some 10 Gyr ago (see e.g. Dinescu et al. 1999; Gnedin et al. 2002; Bekki \\& Norris 2006; Romano et al. 2007; Bellazzini et al. 2008). It is also customarily acknowledged that metal enriched winds are the most straightforward explanation of the observed properties of dwarf galaxies (see the recent review by Tolstoy, Hill \\& Tosi 2009, and references therein).  In the case of $\\omega$\\,Cen's progenitor we have found, indeed, that significant outflow is needed in order to reduce the effective yield per stellar generation and explain the observed properties (the stellar metallicity distribution function, age-metallicity relation and the trends of several abundance ratios with metallicity; Romano et al. 2007; see also Ikuta \\& Arimoto 2000). Following the results of hydrodynamical simulations by Recchi et al. (2001), we have assumed that elements produced by SNe are lost more efficiently than others. The lowest efficiency of ejection through the outflow has been assigned to hydrogen and helium. This is a common choice when dealing with the chemical evolution of local dwarf spheroidals and it leads to a good reproduction of several observed properties (e.g. Lanfranchi \\& Matteucci 2003). However, no direct measurements of He are available for local dwarf spheroidals: since they lack gas, they cannot be observed for optical/radio recombination lines. Thus, the efficiency of He entrainment in the galactic wind is, as a matter of fact, totally unconstrained. With this in mind, the challenging bet is: could $\\omega$\\,Cen's parent galaxy \\emph{retain most of its helium,} while still \\emph{ejecting most of the heaviest species?}  Up to now, three major scenarios have been proposed for the origin of the helium-enriched populations in $\\omega$\\,Cen, as well as in other massive GCs: \\begin{enumerate} \\item[A] accretion of helium-rich material by pre-existing stars; \\item[B] star formation out of the ejecta from either massive AGB stars (B1) or fast rotating massive stars (B2); \\item[C] pollution by Pop\\,{\\sevensize III} stars. \\end{enumerate} Recently, Renzini (2008) has reviewed them critically and concluded that only the AGB option (B1) appears to be acceptable; the others result in either a spread of helium abundances (A and B2) or a certain degree of enrichment in heavy elements as well (C). However, for the AGB option to work, AGB stars more massive than 3~M$_{\\odot}$ must experience just a few TDU episodes and the cluster hosting the He-rich population must be the remnant of a more massive systems. Detailed stellar modelling (Ventura \\& D'Antona 2008a,b) shows that only AGB stars above 4--5~M$_{\\odot}$ can actually reach the required levels of He enhancement in their atmospheres.  In this paper we have thoroughly analysed the issue of the formation of He-rich stars in $\\omega$\\,Cen by means of a complete model for its chemical evolution. In $\\omega$\\,Cen we are facing a \\emph{primary population} (about 50 per cent of the stars, with $\\langle$[Fe/H]$\\rangle \\simeq -$1.7 and `normal' He) and a \\emph{secondary population} (at intermediate, $\\langle$[Fe/H]$\\rangle \\simeq -$1.4, and extreme, $\\langle$[Fe/H]$\\rangle \\simeq -$0.6, metallicities). At least part of the secondary-population stars may be He enhanced.  Although the interaction with the Milky Way is likely to have played some role in shaping the chemical properties of (at least part of) the stellar population in $\\omega$\\,Cen, we find a good agreement between predictions from standard chemical evolution models (that do not take all the relevant dynamical processes into account) and observations of average abundances and abundance ratios by assuming that a galactic-scale outflow vented out mainly matter enriched in SN products, while mostly retaining the ejecta of slow stellar winds, irrespective of the initial mass of the stellar progenitor. Besides this ejection of enriched gaseous matter, the cluster progenitor must have lost a major fraction of stars somehow later on during its evolution. In summary, we show that, in order to explain the existence of extreme He-rich stellar populations, \\emph{what really matters is the kind of evolution the host system went through, rather than the kind of stars responsible for a `super He production'.}  In $\\omega$\\,Cen, because of the relatively long-lasting star formation (some 10$^9$ years), stars of initial mass as low as 2~M$_{\\odot}$ had the time to contribute significantly to the chemical enrichment of the ISM. This is clearly witnessed by the high \\emph{s-}process abundances displayed by some cluster stars (e.g. Johnson et al. 2009). If cooling flows developed to collect the AGB ejecta in the innermost galactic regions where chemically peculiar stars were born, it can be hardly told which selective mechanism brought only the ejecta from 4--5 to 8~M$_\\odot$ AGB stars to the cluster core, while leaving behind those from lower-mass stars. Including a contribution to the chemical enrichment from super-AGB stars could help the AGB star pollution scenario to produce results in agreement with the observations (e.g. Pumo et al. 2008), but up to now no nucleosynthetic yields from this class of objects have been provided in the literature for use in chemical evolution studies.  The competing scenario of pollution from slow winds of FRMSs, while being able to reproduce, in principle, the peculiar abundances and abundance ratios required by the observations of chemically peculiar stars, is hampered by a number of assumptions about, for instance, the efficiency of meridional circulation to mix the products of core H-burning in the stellar envelope, or the rate of mass loss through the mechanical wind, or the need for the development of a large contingent of fast rotators in GCs. Since the matter expelled through the equatorial discs of FRMSs is available with the `right' He abundance since the very beginning of the proto-$\\omega$\\,Cen formation (see Sect.~\\ref{sec:heMASS} and Fig.~\\ref{fig:epMASS}), it is not clear why an extreme He-rich population should form only later on, at metallicities around [Fe/H]~$\\sim -$1.2 (and possibly higher).  As a last comment, it is worth stressing that abundances of He as high as that required to explain the blue MS of $\\omega$\\,Cen have never been observed elsewhere. Although we may think about mechanisms able to originate extreme He-rich stellar populations, we must also be aware that the existence of such extremely high He abundances still await a confirmation from spectroscopic observations."
    },
    "0910/0910.1066_arXiv.txt": {
        "abstract": "The misalignment mechanism for axion production depends on the temperature-dependent axion mass. The latter has recently been determined within the interacting instanton liquid model (IILM), and provides for the first time a well-motivated axion mass for all temperatures. We reexamine the constraints placed on the axion parameter space in the light of this new mass function. Taking this mass at face value, we find an accurate and updated constraint $ f_a \\le 2.8(\\pm2)\\times 10^{11}\\units{GeV}$ or $m_a \\ge 21(\\pm2) \\units{\\mu eV}$ from the misalignment mechanism in the classic axion window (thermal scenario). However, this is superseded by axion string radiation which leads to $ f_a \\lesssim 3.2^{+4}_{-2} \\times 10^{10} \\units{GeV}$ or $m_a \\gtrsim 0.20 ^{+0.2}_{-0.1} \\units{meV}$. In this analysis, we take care to precisely compute the effective degrees of freedom and, to fill a gap in the literature, we present accurate fitting formulas. We solve the evolution equations exactly, and find that analytic results used to date generally underestimate the full numerical solution by a factor $2-3$. In the inflationary scenario, axions induce isocurvature fluctuations and constrain the allowed inflationary scale $H_I$. Taking anharmonic effects into account, we show that these bounds are actually weaker than previously computed. Considering the fine-tuning issue of the misalignment angle in the whole of the anthropic window, we derive new bounds which open up the inflationary window near $\\theta_a \\to \\pi$. In particular, we find that inflationary dark matter axions can have masses as high as 0.01--1$\\units{meV}$, covering the whole thermal axion range, with values of $H_I$ up to $10^9$GeV. Quantum fluctuations during inflation exclude dominant dark matter axions with masses above $m_a\\lesssim 1$meV. ",
        "introduction": "\\label{sec:introduction}  Axions are still one of the best motivated cold dark matter candidates. Initially invented to solve the strong CP problem (``why is the QCD vacuum angle so so small?'', i.e.\\ $\\theta<10^{-9}$), it was soon realised by Weinberg \\cite{weinberg:axion} and Wilczek \\cite{wilczek:axion} that the Peccei-Quinn (PQ) mechanism \\cite{peccei:quinn:cp1,peccei:quinn:cp2} gave rise to a very-light pseudo-scalar Goldstone boson. In order to retain renormalisability, Peccei and Quinn introduced a new chiral symmetry, $U(1)_{PQ}$, on the quark and Higgs fields, that is spontaneously broken. This implies the existence of a new particle, a would-be pseudo-Goldstone boson, the axion; it receives a mass due to instantons because $U(1)_{PQ}$ is anomalous. In the original papers, the axion was incorporated in the electroweak sector but laboratory experiments soon ruled out such a light boson with $\\units{GeV}$ coupling. This gave rise to the so-called invisible axion models \\cite{shifman:vainshtein:zakharov:cp,kim:axion,dine:fischler:srednicki:axion,zhitnitsky:axion}, that a priori are not tied to any known energy scale. To constrain them, it was realised that such extremely weakly interacting particles could provide a new cooling mechanism for stars. The invisible axions have typically very weak couplings to ordinary matter. On the one hand, this makes their experimental detection difficult but, on the other hand, provides us with a well-motivated dark matter candidate. Refer to past reviews \\cite{kim:report,cheng:strongCP:report,turner:axion:report} for further details.  The axion has a rich phenomenology in that it can be produced thermally or non-thermally. The thermal production channel is the standard scenario for most WIMP's \\cite{turner:kolb:cosmology}. Recently it was shown that the thermal axion cannot contribute the dominant dark matter component of the universe \\cite{hannestad:mirizzi:raffelt:thermal:axion}\\footnote{The thermal axion bound follows from $(m_a^{th}/130\\units{eV}) < \\Omega_c\\, g^{dec}_{*}/10$, and is saturated for $m_a^{th} \\approx 15 \\units{eV}$. This bound is, however, excluded by the new astrophysics bound $m_a<0.01\\units{eV}$. Thus, $\\Omega_a^{th}=\\Omega_c(0.01/15)\\approx 0.001 \\,\\Omega_c$.}. The axion can also be produced non-thermally: after the spontaneous breaking of the PQ symmetry, the axion lives in a $U(1)$ vacuum manifold; such a broken field supports the formation of topological strings \\cite{davis:axion:string:1985,davis:axion:string:1986,davis:shellard:axion:string,dabholkar:quashnock:axion:string,battye:shellard:axion:string:1994b} \\cite{harari:sikivie:axion:string,hagmann:sikivie:axion:string,hagmann:chang:sikivie:axion:string}, whose radiation produce axions. Finally, axions can be produced non-thermally through the so-called misalignment mechanism: at the QCD phase transition non-perturbative effects generate a mass, and the axion field relaxes to its minimum, which is precisely the PQ mechanism\\footnote{The $\\theta$ angle, a free parameter, is replaced by a dynamical field that evolves to its CP-conserving minimum.}, invented to solve the strong CP problem. The oscillation around its minimum produce a coherent state of zero mode axions, i.e.\\ a Bose-Einstein condensate \\cite{sikivie:yang:axion}. This last production scenario is potentially sensitive to the QCD effects, i.e.\\ the axion mass, and is the primary subject of this paper.  Because of the anomalous $U_\\mathrm{PQ}(1)$ symmetry, the axion has a two gauge boson interaction and can thus decay into two photons; such processes are used to look for axions experimentally, e.g.\\ in solar axion searches and vacuum birefringence experiments. The former is one of the more stringent astrophysical constraints, the strongest coming from the analysis of the supernova 1987A neutrino flux which would be affected by axions. It gives a lower bound for the axion decay constant, $f_a \\gtrsim 10^9 \\mathrm{GeV}$. See \\cite{kuster:raffelt:beltran:axions} for a recent, comprehensive set of review articles.  In section \\ref{sec:CP:PQ} we briefly review the strong CP problem and the PQ mechanism that lead to the introduction of the axion. We continue to discuss the effective axion potential in section \\ref{sec:axion:potential}, and essentially focus on the temperature dependent axion mass. We review our determination of the mass in the framework of the interacting instanton liquid model (IILM) \\cite{wantz:iilm:3}. In section \\ref{sec:axion:cosmology} we reexamine the cosmology of the vacuum realignment production mode in light of this new mass function: We solve the cosmological evolution equations numerically, and compare the results to the standard analytic approximation, identifying regimes in which present estimates are and are not robust. To make the numerics self-contained, we include the correct effective degrees of freedom for the entropy and radiation density and provide accurate fitting formulas for their rapid evaluation in appendix \\ref{app:geff:fit}. For completeness, we also review and update constraints from axion string radiation. ",
        "conclusions": "We have presented a temperature-dependent axion mass, based on instanton methods, that is valid for all temperatures. The transition between the high- and low-temperature regime is well-motivated and computed within the same model, the IILM, in contrast to some ad hoc procedure to connect the two.  Although the IILM does not explain confinement, chiral symmetry restoration is incorporated and the model can be expected to give qualitatively correct results for the axion mass; this relies on the fact that it is related to the QCD topological susceptibility which in turn is a chiral quantity. We note that chiral symmetry restoration is indeed seen in the IILM, although at a slightly lower temperature. Given the discovery of the more general KvBLL calorons, that may play an important role in the confinement/deconfinement transition, we expect to improve our understanding of the axion potential in the future by incorporating these new degrees of freedom into the IILM.  Using this new axion mass, we solved numerically the axion evolution equations in the concordance FRW cosmology. It turns out that the analytic approximations used previously differ by a factor of $2-3$. This is unexpectedly good agreement, considering the crude determination of the axion mass within the dilute gas approximation. We believe it to be the result of a coincidence, that the extrapolation of the high temperature DGA axion mass fairly closely follows the full IILM result around the phase transition. Conversely, this correspondence can be interpreted as evidence that the axion mass determination is fairly robust. This is also seen from the rather small differences between the numerical results between the IILM mass and the lattice inspired mass.  We want to draw to attention that the IILM is a model of QCD, and as such it needs to be checked against lattice data, say. In that light, there remains the possibility that the true high temperature axion mass is different from the IILM prediction. This could lead to qualitatively different results: On the one hand, a softer decay, as seen in gluodynamics, will lead to weaker constraints. A more abrupt mass switch on, on the other hand, would tighten the constraints and potentially close the classic axion window. The ideas of molecule formation within the IILM, and the subsequent stronger suppression of the axion mass in the plasma phase, have initially prompted this investigation. Within the current IILM this is not realised, but it is not ruled out either. The lattice community is performing realistic QCD simulations directly at the physical quark masses, and we can expect a reliable axion mass determination to follow from that data in future.  To get an accurate estimate across the whole $f_a$ axis, we included the correct temperature dependence between the scale factor and the plasma temperature which follows from the conservation of entropy. To that end we computed the full phase-space integral to get the temperature dependence for the effective degrees of freedom $g_{*.S}$ and $g_{*,R}$, following closely \\cite{coleman:roos:geff}. Additionally, we provide accurate fitting formulas, which to our knowledge have not been presented in the literature previously.  In the classic axion window, where the PQ symmetry breaks only after inflation, a quantitative analysis of the misalignment bound yields $f_a < 2.8(\\pm2) \\times 10^{11}\\units{GeV}$, $m_a>21(\\pm 2) \\units{\\mu eV}$. While the misalignment axion is a useful benchmark for comparing alternative estimates, it is far exceeded in density by axions from string radiation which yields the more stringent string bound $f_a~\\lesssim~3.2^{+4}_{-2} \\times 10^{10} \\units{GeV}$. Experimental searches for a thermal dark matter axion should focus on masses around $m_a \\sim 200 \\units{\\mu eV}$.  The anthropic axion window, which entangles $\\theta_a$ and $f_a$, is constrained by investigating the production of isocurvature perturbations. The rather small contribution of these to the CMBR power-spectrum places strong bounds on $f_a$ and the inflationary scale $H_I$. This bound is strong enough so that quantum fluctuations are totally negligible, i.e.\\ $\\sigma_a \\ll \\langle \\theta_a \\rangle$, for almost all axion angles. The anthropic window allows a dark matter axion that can have any initial misalignment angle, provided the inflationary scale is low enough. In light of naturalness in fundamental theory, most interest has focused on the constraints for large $f_a$.  Nevertheless, taking the whole of the anthropic range seriously reveals some interesting possibilities. In particular, we investigated the region with large misalignment angles, especially $\\theta_a \\to \\pi$. In this regime the axion potential can no longer be described by a parabola, and we need to take the anharmonic effects into account. A self-consistent solution follows for the regime with small fluctuations, in which case we can Taylor-expand the axion potential around $\\langle \\theta_a \\rangle$; the anharmonic effects are then encoded in the anharmonic dependence on $\\langle \\theta_a \\rangle$. These effects lead to smaller isocurvature perturbations as compared to the harmonic case and, consequently, lead to weaker constraints on the inflationary scale.  Fine-tuning the initial misalignment angle to $\\pi$, we find that the quantum fluctuations must be even more stringently constrained. If the axion provides the dominant component to the dark matter content of the universe and if we take at face value an inflationary period after the spontaneous breaking of the PQ symmetry, then this places the strongest constraint on $f_a$ and $H_I$ at large initial misalignment angle. We note the intriguing possibility of a dominant dark matter axion with a mass $m_a \\sim 200 \\units{\\mu eV}$ which is consistent with either the thermal or the inflationary scenarios. This model independence provides extra motivation for experimental searches around this mass range. Anthropic tuning near $\\theta\\approx \\pi$ allows inflationary dark matter axions ($\\Omega_a = 0.23$) to have masses as high as $m_a\\le 1 \\units{meV}$, but no higher because quantum fluctuations restrict the fine tuning (for $H_I \\ge 10^4 \\units{GeV}$).  We note that the isocurvature and quantum fluctuation constraints become weaker if the axion is not the dominant dark matter candidate, and we displayed the ensuing 3-dimensional parameter space.  Finally, the dependence of the cosmological constraints on the axion mass is only a secondary effect for the isocurvature constraint and the axion string contribution. In the former case the examined parameter space is so large that only significant modifications play a role, such as the change from a harmonic to an anharmonic potental; in the latter case, the uncertainties inherent in the computation easily exceed those from the axion mass dependence."
    },
    "0910/0910.1591_arXiv.txt": {
        "abstract": "Stellar masses play a crucial role in the exploration of galaxy properties and the evolution of the galaxy population. In this paper, we explore the minimum possible uncertainties in stellar mass-to-light (M$_\\ast$/L) ratios from the assumed star formation history (SFH) and metallicity distribution, with the goals of providing a minimum set of requirements for observational studies.  We use a large Monte Carlo library of SFHs to study as a function of galaxy spectral type and signal-to-noise ratio (S/N) the statistical uncertainties of M$_\\ast$/L values using either absorption-line data or broad band colors.  The accuracy of M$_\\ast$/L estimates can be significantly improved by using metal-sensitive indices in combination with age-sensitive indices, in particular for galaxies with intermediate-age or young stellar populations. While M$_\\ast$/L accuracy clearly depends on the spectral S/N ratio, there is no significant gain in improving the S/N much above 50/pix and limiting uncertainties of $\\sim 0.03$~dex are reached.  Assuming that dust is accurately corrected or absent and that the redshift is known, color-based M$_\\ast$/L estimates are only slightly more uncertain than spectroscopic estimates (at comparable spectroscopic and photometric quality), but are more easily affected by systematic biases.  This is the case in particular for galaxies with bursty SFHs (high H$\\delta_A$ at fixed D4000$_{\\rm n}$), the M$_\\ast$/L of which cannot be constrained any better than $\\sim0.15$~dex with any indicators explored here. Finally, we explore the effects of the assumed prior distribution in SFHs and metallicity, finding them to be higher for color-based estimates. ",
        "introduction": "\\label{sec:intro} The study of the galaxy population as a function of their stellar masses has become increasingly common and important in the last years. For a long time, galaxy scaling relations were studied in terms of galaxy luminosities, often in the optical regime where differences in star formation histories between different galaxies would lead to factors of 3-10 difference in luminosity for a given stellar mass. Comparison of galaxies with different star formation histories and similar luminosities is therefore not a well-posed exercise, because it translates into comparing galaxies of dramatically different stellar mass.  This situation has significantly improved in the last decade thanks to the development and continuous refinement of methods to estimate the stellar mass-to-light ratio ($\\rm M_\\ast/L$) from galaxies spectro-photometric properties. All the various techniques involve comparison with predictions from population synthesis models. The observational constraints used to derive galaxy $\\rm M_\\ast/L$ estimates range from relatively `cheap' observations of colors \\citep[e.g.][]{BE00,papovich01,BdJ01,cole01,bell03}, to broad-band spectral energy distributions (SEDs) extending from the optical or UV up to the IR \\citep[e.g.][]{borch06,walcher08,franzetti08}, to spectroscopic measurements of individual stellar absorption features \\citep[e.g.][]{kauffmann03,gallazzi05} or the full optical galaxy spectrum \\citep[e.g.][]{panter04,tojeiro07,cidfernandes07}.  The technical development has proceeded in parallel with (and has been largely motivated and made possible by) the rise of large homogenous spectroscopic and photometric datasets from surveys such as the Sloan Digital Sky Survey \\citep{york00}, the Two Micron All-Sky Survey \\citep{2mass}, COMBO-17 survey \\citep{wolf04}, VIMOS VLT Deep Survey \\citep{lefevre05}, and COSMOS \\citep{scoville07}. The application of stellar mass estimation techniques to these datasets has allowed the community to make rapid progress towards understanding the properties and evolution of the galaxy population \\citep[e.g.,][]{kauffmann03b,shen03,baldry06,bell03,fontana06}.  While it is generally acknowledged that ranking galaxies by their estimated stellar masses rather than observed luminosities provides a more physical insight, it is also recognized that in translating observational quantities into physical parameters estimates, we are limited by a number of statistical and systematic uncertainties. To start with, the accuracy of physical parameters  estimates clearly relies on the accuracy of the population synthesis models used to interpret observations (i.e. galaxy colors or spectra). In particular differences in the stellar libraries (either empirical or theoretical) and uncertain (not-well observationally constrained and/or not-well theoretically understood) aspects of particular stellar evolutionary phases can affect the predicted SED for a given star formation history (SFH). Among these, it is recognized that thermally-pulsating asymptotic giant branch stars (TP-AGB) can have a significant impact on $\\rm M_\\ast/L$ values estimated from near-infrared (NIR) colors of young stellar populations ($\\sim 0.5-2$Gyr). For instance, $\\rm M_\\ast/L$ of high-redshift galaxies based on \\cite{maraston05} models can be lower by 40-60\\% than those based on \\cite{bc03} code because of the different flux contribution by stars in this particular phase \\citep[e.g.,][]{maraston06,bruzual07,cimatti08}. Blue stragglers and blue horizontal branch morphologies can instead affect (in an age-dependent fashion) the SEDs, and derived properties, of passive galaxies. Moreover, the treatment of non-solar abundance ratios in population synthesis models can affect the stellar ages (and hence $\\rm M_\\ast/L$) of early-type galaxies derived from spectral features \\citep[e.g.,][]{thomas04}. These and other stellar populations uncertainties have been extensively explored in \\cite{conroy08,conroy09}.  Dust attenuation is a key uncertainty in stellar M/L values, especially those derived using color information (e.g., \\citealt{pozzetti07}, \\citealt{fontana04}, \\citealt{zibetti09} - although dust can affect absorption features like the 4000\\AA-break, \\citealt{macarthur05}). It has been argued that dust should to first order cancel out in the total M$_\\ast$ estimated \\citep{BdJ01}. Yet, there are a number of caveats and possible exceptions to this simplistic picture.  In particular, for systems with high optical depth (either globally or in patches), there can be considerable absorption of light without large amounts of reddening \\citep{deJong96,driver07,zibetti09}.  If the attenuation is patchy, moving towards resolved mass maps, and then integrating the result to a total mass, reduces this systematic uncertainty \\citep{abraham99,zibetti09}\\footnote{This is an advantage of color-based stellar M/L estimates over spectroscopic values: the spectra to be analyzed typically are only of a part of a galaxy, whereas imaging exists across the face of the galaxy and one can account for optically-thick parts separately, if one wishes.}.  Modeling of the effects of dust can help \\citep{deJong96,driver07},     where total thermal infrared flux can be of some use in constraining the model \\citep[e.g.,][]{popescu05,dacunha08}.  The predicted photometric properties and $\\rm M_\\ast/L$ of galaxies also strongly depend on the assumed initial mass function (IMF). While it is sometimes possible to convert $\\rm M_\\ast/L$ values determined assuming a given IMF into another one with a simple scaling factor when the differences are only at masses $\\la 1 M_{\\sun}$, IMFs that are different at masses $\\ga 1 M_{\\sun}$ can alter the galaxy SED in an age-dependent way and hence have a different global effect on the colors of star-forming and passive galaxies \\citep{conroy09}. Generally, studies of the average slope of the high-mass ($>1$~M$_\\odot$) part of the stellar IMF determine slopes consistent with Salpeter \\citep[e.g.][]{kennicutt83,BG03,hoversten08}. There are a number of studies that rely on chemical abundances that sometimes favor slightly top-heavier IMFs \\citep[e.g.][]{worthey92,gibson97,thomas99,trager00,arrigoni09}; yet, note that chemical modeling of galaxy evolution is uncertain (owing to uncertainties in e.g., yields, gas infall and outflow) and such results should be viewed as being interesting but necessarily tentative in nature.  Moreover it is generally assumed that the shape of the IMF is universal and constant in time. Dynamical and strong lensing estimates of total M/L on galactic scales  (i.e., $\\la$ half-light radius scales) have historically proven  compatible with a universally-applicable \\citet{chabrier03} or  \\citet{kroupa01} stellar IMF \\citep[e.g.,][R.\\ S.\\ de Jong \\& E.\\ F.\\ Bell, in preparation]{BdJ01,cappellari06,gallazzi06,ferreras08}; such an IMF is  motivated by local observations of the luminosity/mass function of  star clusters, and is consistent with the statistics of  H$\\alpha$ equivalent width as a function of galaxy color for luminous galaxies \\citep{hoversten08}.  There have been a few recent arguments that, if borne out by further analysis and observations, may suggest a time-dependent stellar IMF \\citep{wilkins08,dave08,vandokkum08}. Furthermore, statistics of H$\\alpha$ equivalent width as a  function of galaxy color/UV flux show a tendency towards fewer very massive stars in dwarf galaxies \\citep{hoversten08,meurer09}  that may indicate a change in the galaxy-averaged stellar IMF (as opposed to the IMF of individual clusters) in low-luminosity galaxies (see also, e.g., \\citealp{weidner05}). It has to be noted however that \\cite{elmegreen06} found no evidence of the composite IMF being steeper that individual cluster IMF, supporting the equality of the IMF from small to large scales.  Finally, we note that the derived $\\rm M_\\ast/L$ estimates depend on the assumed distribution in SFHs of the models used to interpret galaxy SEDs. In particular it has been shown that the addition of bursts of star formation on top of a continuous SFH can produce $\\rm M_\\ast/L$ estimates systematically different by an amount that can vary between $\\sim$10\\% and a factor of two depending on the strength and fraction of starbursts \\citep[e.g.,][]{BdJ01,drory04,pozzetti07,gallazzi08,wuyts09}. In the last few years several works have put significant efforts in quantifying the overall statistical and systematic uncertainties on $\\rm M_\\ast/L$ estimates and quantities based on them, such as stellar mass functions and cosmic stellar mass densities \\citep[to mention a few][]{cimatti08,gallazzi08,marchesini09,LS09}. The general impression from all such works is that it is difficult to defend $\\rm M_\\ast/L$ estimates to much better than 0.1~dex under the most ideal conditions, and perhaps even worse if more extreme conditions are allowed as discussed above (e.g. if there are systematic and/or random stellar IMF variations from galaxy to galaxy, if there are frequent prominent bursts).  In this paper, we take a somewhat orthogonal approach to the question of $\\rm M_\\ast/L$ accuracies. In particular we neglect here uncertainties related to the physics of population synthesis models and assume that we have a `perfect' model and that the IMF is indeed universally-applicable. We further assume that dust can be ignored, either because dealing with dust-free systems or because it can be `perfectly' corrected for. The motivation for making such ideal assumptions is to isolate the contribution of SFH and metallicity scatter to the $\\rm M_\\ast/L$ error budget. By doing so we wish to explore in depth the following issues: i) how does $\\rm M_\\ast/L$ accuracy depend on the choice of stellar population diagnostics used to estimate its value; ii) how does $\\rm M_\\ast/L$ accuracy respond to increasing spectroscopic or photometric quality; iii) how do the answers to these questions depend on the galaxy SFH (or spectral type).  While at first sight our approach may seem somewhat notional, given the aforementioned uncertainties in stellar population models and concerns about the possibility of stellar IMF variations, yet in isolating the influence of SFH and metallicity variations on $\\rm M_\\ast/L$ values one can gain insight into a number of issues. \\begin{itemize} \\item How much effort should be placed in creating realistic and physically-motivated grids of SFH and metallicity for use in stellar mass estimation?  If variations in SFH and metallicity produce little scatter in $\\rm M_\\ast/L$ values, then little effort needs to be expended in producing realistic templates.  If SFH/metallicity variations lead to large scatter in stellar M/L, even for excellent data, then the creation of realistic priors for the SFH and metallicity distributions acquires some urgency. \\item When should we stop trying to improve the models?  For example, if a particular source of uncertainty in a stellar population model gives an improvement that is much smaller than the minimum possible uncertainty from plausible variations in $\\rm M_\\ast/L$ from SFH/metallicities, then one could argue that that particular improvement in stellar population model is less urgent than other sources of uncertainty. \\item How good does my data need to be?  An investigation of the scatter in  $\\rm M_\\ast/L$ from SFH/metallicity variations at a given set of colors or line indices, for a given data quality, gives a lower limit to the $\\rm M_\\ast/L$ uncertainties achievable with such data. Given a target $\\rm M_\\ast/L$ uncertainty and method, one then learns what the target data quality should be (for effects related to SFH/metallicities only).  This is less of a driver of this work, as higher quality or complementary data may be necessary to address other sources of uncertainty (such as stellar population synthesis model uncertainties or IMF variations); nonetheless, this work helps to build intuition as to what kind of data quality might be required to produce accurate $\\rm M_\\ast/L$ estimates, assuming that the other sources of uncertainty are dealt with adequately. \\end{itemize}  In order to address these issues we study in some detail a large Monte Carlo library of SFHs described in Section~\\ref{sec:library} that cover the spectroscopic and photometric properties of present-day galaxies. We randomly select a subset of models that represent different galaxy spectral types having different SFHs, and we create mock galaxy samples by perturbing their spectra according to a given signal-to-noise ratio (S/N).  We follow a Bayesian approach to derive the likelihood distribution of M$_\\ast$/L for these mock galaxies adopting as observational constraints several sets of absorption indices as outlined in Section~\\ref{sec:estimates}. In Section~\\ref{sec:results} we discuss the dependence of spectroscopically-derived $\\rm M_\\ast/L$ estimates on the absorption indices used, on the spectroscopic S/N, on the SFH scatter, both in general and as a function of galaxy spectral type. We then compare these results to those concerning $\\rm M_\\ast/L$ estimates based on optical or optical-NIR colors in Section~\\ref{sec:colors}. In Section~\\ref{sec:prior} we explore the effects on $\\rm M_\\ast/L$ estimates of any mismatch between the assumed SFH and metallicity prior and the true distribution. We finally present our conclusions in Section~\\ref{sec:conclusions}. Throughout this work we adopt a \\cite{chabrier03} IMF with mass cut-offs of 0.1 M$_\\odot$ and 100 M$_\\odot$ and we refer to the $z$-band stellar mass-to-light ratio. ",
        "conclusions": "\\label{sec:conclusions}  \\subsection{Recap}  In this work we have explored the {\\it statistical} uncertainties, isolating the effects of SFH and metallicity variations, affecting stellar mass-to-light ratio estimates derived by comparing either stellar absorption features or colors with predictions from a comprehensive library of SFHs. In particular we have addressed the question of which spectroscopic constraints and which spectral accuracy (as quantified by the S/N) are necessary/helpful to reduce the uncertainties as a function of galaxy type.  We made use of recent medium-resolution population synthesis models \\citep{bc03} to generate spectral absorption indices and broad-band colors for a Monte Carlo library of SFHs characterized by both exponentially declining and bursty SFHs. We randomly selected $\\sim550$ model galaxies in order to cover different positions in the \\dn-\\hd\\ plane which represent different average SFHs: continuous SFHs following a sequence of increasing \\hd\\ (and decreasing stellar age) at decreasing \\dn, and bursty SFHs identified by higher \\hd\\ values at fixed \\dn. We then generated samples of mock galaxies at different S/N by perturbing the absorption indices of the selected models according to the typical observational error at that S/N. For each mock galaxy we estimated the likelihood distribution of its $z$-band mass-to-light ratio (\\mtol) by comparing different sets of absorption indices with those of all the models in the library. We quantify the accuracy on M$_\\ast$/L estimates by both the median offset from the true M$_\\ast$/L and the median 68\\% likelihood interval of M$_\\ast$/L computed from the PDF.  We have considered different combinations of absorption indices by including age-sensitive indices only (\\dn\\ and Balmer lines), adding metal-sensitive indices (either extending to red wavelengths in order to include Mg and Fe lines, or only in the blue portion of the spectrum $\\la5000$\\AA), and using only the red portion of the optical spectrum (in particular excluding \\dn).  The uncertainty on \\mtol\\ estimates shows a similar dependence on S/N regardless of the spectral features used, decreasing significantly with S/N up to a S/N of 30 and then only mildly.  In particular, for the sample as a whole, a S/N per pixel of at least $\\sim30$ is required to constrain \\mtol\\ to within better than 30\\% when using age-sensitive indices only (\\dn\\ and Balmer lines). However there is no significant gain by increasing further the S/N: there seems to be a limiting uncertainty of 0.07~dex, when using \\dn\\ and \\hd\\ only, or 0.05~dex when using all the three Balmer lines in addition to \\dn. At given S/N the best results are obtained by combining \\dn\\ and Balmer lines with Mg and Fe features on a relatively long wavelength baseline, which constrain not only the age of the stellar populations but also their metallicity (and thus narrow the allowed \\mtol\\ range). In this case, errors below 0.05~dex could be reached with S/N$\\ga50$ spectra. When only spectra below $\\sim5000$\\AA~ are available, there is no significant gain in using additional absorption features with respect to \\dn\\ and \\hd\\ only. On the other hand, when only the red portion of the spectrum is available (hence no measure of \\dn\\ is possible) similar results on \\mtol\\ are obtained by combining \\hb\\ with metal-sensitive features.  By distinguishing galaxy spectral type, we find in general that \\mtol\\ is easier to constrain (i.e. lower S/N and fewer observational constraints are required) for old stellar populations: for example, uncertainties below $\\pm0.05$dex can be reached at S/N$\\sim20$ with \\dn\\ and \\hd\\ only. When using \\dn\\ and \\hd\\ only, \\sigmamtol\\ depends on stellar age in the sense that it becomes larger for galaxies with younger luminosity-weighted ages. For such galaxies, in particular those dominated by a recent burst of star formation, which have a broader metallicity distribution, it is crucial to add constraints from metal-sensitive lines in order to reach accuracy $\\la0.1$dex at reasonable S/N.  The stellar mass-to-light ratio of galaxies at intermediate \\dn\\ and high \\hd\\ values, characterized by recent bursts of SF and intermediate-age populations, cannot be constrained from spectroscopy within better than $0.15-0.2$dex. The ability of reducing \\sigmamtol\\ for galaxies with bursty SFHs (in general with high \\hd\\ values) is also limited by discreteness effects which come into play at S/N$\\ga50$ and using more than two observational constraints, in which cases the PDF cannot be reliably sampled. Even if there were enough models to build a robust PDF, the intrinsic scatter in \\mtol\\ at fixed index-index values is between 0.1 and 0.2~dex at intermediate \\dn\\ and high \\hd. We conclude that for this type of galaxies there is no gain in increasing the S/N above $\\sim30$ or using more observational constraints.  Applying the same Bayesian method to derive \\mtol\\ estimates from colors shows that, in the assumption that redshift is known and neglecting dust effects, it is possible to reach an accuracy only slightly worse than with absorption indices. However a photometric accuracy of at least 0.05~mag is required. The smallest uncertainties are obtained with $g-i$ which has the widest wavelength leverage. As opposed to optical colors, optical-NIR colors of similar photometric quality provide larger uncertainties because they show a broader relation with \\mtol\\ and they depend on both age and metallicity. Moreover they tend to under-/over-estimate the true \\mtol\\ for old/young populations, respectively.  We note that, while (optical) colors provide similar or smaller (and more homogeneous) uncertainties for bursty, intermediate-age galaxies, they systematically overestimate their \\mtol\\ by $\\sim0.2$dex. There is no significant gain in \\mtol\\ accuracy by adding another color to $g-i$. Although uncertainties are slightly smaller by combining $g-i$ with a NIR color (such as $i-H$), this is not enough to remove the bias for bursty, intermediate-age galaxies.  \\subsection{The Bottom Line}  We conclude that: \\begin{itemize} \\item The mass-to-light ratio of galaxies dominated by old stellar populations or with smooth SFH is in general better constrained by spectroscopic absorption indices (in particular combining age- and metal-sensitive features for young stellar populations) both in terms of retrieving the correct value and in terms of statistical uncertainty. An accuracy below 0.1~dex is reached already at S/N$\\ga20$ and it can go below 0.05~dex at higher S/N. \\item For galaxies with smooth SFH and young stellar populations  one optical color or a combination of an optical and an optical-NIR color provides only slightly worse statistical uncertainties on \\mtol\\ than spectroscopic constraints, in the assumption that redshift is known and dust is not important, and when photometric errors are below $\\sim0.05$mag. \\item For galaxies which have experienced a recent burst of SF (identified by high \\hd\\ values at fixed \\dn) we are limited by either large uncertainties ($\\sim0.15$dex) on spectroscopic estimates or biased estimates from colors (overestimated by up to 0.2~dex). \\end{itemize}  Another important conclusion is that mismatches in the assumed SFH and metallicity distribution can have an impact on the estimated M$_\\ast$/L. It is necessary to include burst of SF in order to reproduce the whole range of observables. However, over-representing them can lead to underestimate the M$_\\ast$/L of old stellar populations by $\\la0.05$~dex if based on absorption features or up to 0.1~dex if based on colors. Perhaps surprisingly, the assumption of constant metallicity does not affect the M$_\\ast$/L estimates from spectra, while it may introduce significant biases in color-based estimates.  \\subsection{Recommendations and Practical Considerations}  We finish by highlighting a few practical issues, with some emphasis on responses to the particular questions posed in the introduction to this paper.  One of the prime motivations of this paper was to understand to which extent data quality affects the determination of M$_\\ast$/L. If high-quality data allowed one to uniquely select a best-fit model out of a sample of more than a hundred thousand SFHs, then relatively little thought needs to be put into the construction of such a model template set for M$_\\ast$/Ls (in particular in the exact distribution of different SFHs and metallicities). Unfortunately, the results of this paper have argued against this perspective. In at least few cases, the stellar M/L values one gets depend sensitively on the exact mix of SFHs in the model template set. The first, perhaps obvious, consideration is that the template set should cover the full range of possible SFHs in order to be applicable to all galaxy spectral types. This becomes critical when good quality data are available. In particular, recent bursts of SF should be well sampled in order to analyse intermediate-age bursty galaxies, where there is considerable bias and scatter in estimated M$_\\ast$/L, even at very high S/N and using a wide range of spectroscopic indices to constrain the model fits (summarized in, e.g., Figs.~\\ref{fig:err_sfh} and~\\ref{fig:col_idx}. On the other hand, the exact fraction of bursts assumed affects the M$_\\ast$/L estimates of old stellar populations (Fig.~\\ref{fig:SFH_spec} and ~\\ref{fig:SFH_col}). Another case is illustrated in Fig.~\\ref{fig:m2l_color_dust}, where it is clear that the time at which star formation is assumed to start makes a considerable difference to (at least) color-derived stellar M/L. In all these cases, any mismatch between the mix of SFHs in the input template set and the actual Universe {\\it at the redshift of interest} will manifest itself through biases and increased scatter in stellar M/L estimates. Our analysis also shows that, while the assumption of single-metallicity stellar populations appears not to affect spectroscopic M$_\\ast$/L estimates (as long as the range in assumed metallicities encompass the range of true mass-weighted metallicities), it can introduce significant biases in color-based estimates. These considerations underline one of the key conclusions of this paper, that the mix of star formation histories and metallicities of the input template set should match as well as possible the actual distribution of star formation histories and metallicities in order to produce optimal stellar M/L values.  A second motivation was to try to understand if and where model development acquires lower priority and urgency, at least from the perspective of stellar mass estimates.  There are parts of parameter space where plausible data quality could yield stellar masses accurate to $\\sim 0.03$\\,dex, in the absence of any systematic mismatches between the model and data.  This potentially very high level of accuracy is of course currently out of reach, given the systematic uncertainties inherent to current stellar population modeling \\citep[e.g.][]{maraston06,conroy08}. Accordingly, it has to be concluded that model development and systematic error control should remain a very high priority for those wishing to understand the stellar mass content of the Universe.  A final motivation was to understand what kind of data is necessary or sufficient  to produce accurate stellar M/L values.   We cannot and do not offer a definitive answer to this question in this paper, for at least two reasons.  Firstly, constraint of important sources of systematic stellar population model uncertainty (e.g., importance of TP-AGB stars to the near-infrared light or contribution of BHB stars to the optical light of old stellar populations) may require additional data not considered here.  Secondly, dust has been neglected here, and further investigation may conclude that more data will be required to model its effect properly than are required simply to constrain the stellar populations.  Notwithstanding these caveats,  we find that, while in general spectroscopic information is superior in constraining a galaxy SFH and hence its M$_\\ast$/L, in the case of relatively smooth star formation histories optical colors alone are enough to constrain the stellar M/L to reasonable accuracy ($0.05-0.1$dex).  The addition of extra broad-band information does not offer large improvements.  An obvious and important case in which the addition of spectral information can improve the quality of M$_\\ast$/L estimates is where there is highly-structured SFH in the last $1-2$Gyr, where Balmer line information can be critical in minimizing the bias in inferred stellar M/L value. Moreover, according to our analysis, spectroscopic M$_\\ast$/L estimates appear preferable to color-based ones in that they are less sensitive to mismatches between the assumed and the true SFH distribution and the assumption of constant metallicity in time."
    },
    "0910/0910.1072_arXiv.txt": {
        "abstract": "Diffusive shock acceleration (DSA) at relativistic shocks is expected to be an important acceleration mechanism in a variety of astrophysical objects including extragalactic jets in active galactic nuclei and gamma ray bursts.   These sources remain strong and interesting candidate sites for the generation of ultra-high energy cosmic rays. In this paper, key predictions of DSA at relativistic shocks that are salient to the issue of cosmic ray ion and electron production are outlined. Results from a Monte Carlo simulation of such diffusive acceleration in test-particle, relativistic, oblique, MHD shocks are presented. Simulation output is described for both large angle and small angle scattering scenarios, and a variety of shock obliquities including superluminal regimes when the de Hoffman-Teller frame does not exist. The distribution function power-law indices compare favorably with results from other techniques.  They are found to depend sensitively on the mean magnetic field orientation in the shock, and the nature of MHD turbulence that propagates along fields in shock environs. An interesting regime of flat spectrum generation is addressed, providing evidence for its origin being due to shock drift acceleration. The impact of these theoretical results on gamma-ray burst and blazar science is outlined. Specifically, {\\it Fermi} gamma-ray observations of these cosmic sources are already providing significant constraints on important environmental quantities for relativistic shocks, namely the frequency of scattering and the level of field turbulence. ",
        "introduction": "\\label{sec:Introduction} There is bountiful evidence for efficient particle acceleration at collisionless shocks in the universe.  The heliosphere, with its planetary bow shocks and traveling interplanetary shocks, has provided interesting and useful test cases for shock acceleration theories.  In the remote regions of the universe, supersonic jets from active galactic nuclei (AGNs) and gamma-ray bursts (GRBs) have offered fascinating windows into an energetic part of the cosmos, where non-thermal radio waves, X-rays and gamma-rays abound.  Fully understanding these sources mandates knowledge of how the particles that generate their light emission are energized.  The foremost paradigms invoke acceleration at relativistic shocks.  Radio imaging reveals very structured and time-variable jets in AGNs, rapid X-ray and gamma-ray variability in both blazar AGNs and GRBs suggest compact emission regions with relativistic bulk motions.  Accordingly, comprehending the relationship between shock acceleration predictions and observations of these sources offers the key to elucidating the understanding of their environs; this constitutes the focus of this paper.  To effect this, an investigation of the features of diffusive shock acceleration is presented, using results from a test particle Monte Carlo simulation.  Then the paper addresses probes of the acceleration theory parameter space imposed by extant GRB and blazar observations in high energy gamma-rays.  The Monte Carlo approach \\cite{EJR90,ED04,NO04,SBS07} is one of several major techniques devised to model particle acceleration at relativistic shocks; others include semi-analytic solutions of the diffusion-convection equation \\cite{KS87,KH89,Kirk00}, and particle-in-cell (PIC) full plasma simulations \\cite{Hoshino92,Nishikawa05,Medvedev05,Spitkovsky08}. Each has its merits and limitations.  Tractability of the analytic approaches generally restricts their solutions to power-law regimes for the phase space distributions \\teq{f(\\hbox{\\bf p})}.  PIC codes are rich in their information on shock-layer electrodynamics and turbulence.  However, to interface with astrophysical spectral data, a broad dynamic range in momenta is desirable, and this is the natural niche of Monte Carlo simulation techniques. A core property of acceleration at the relativistic shocks is that the distribution functions \\teq{f(\\hbox{\\bf p})} are inherently anisotropic. This renders the power-law indices and other distribution characteristics sensitive to directional influences, such as the magnetic field orientation with respect to the shock normal, and the nature of MHD turbulence that often propagates along the field lines.  These connections between observables and physical parameters of the shock environs are studied in some depth in this exposition. ",
        "conclusions": "\\label{sec:conclusion}  This paper has investigated some of the key characteristics of particle acceleration at relativistic shocks, including the identification of the role of shock drift acceleration in generating flat distributions in mildly-relativistic, subluminal shocks.   It has also explored the connection between the acceleration process and high energy emission in two classes of astrophysical sources, namely gamma-ray bursts and jets in blazars. The simulation results presented clearly highlight the non-universality of the index of energetic, non-thermal electrons and ions, spawned by the variety of shock obliquities and the character of hydromagnetic turbulence in their environs.  This non-universality poses no problem for modeling GRB or blazar high-energy power-law indices, though observations generally constrain the parameter space to subluminal or highly-turbulent and modestly superluminal shocks not far from the Bohm diffusion limit.  Diffusive acceleration at ultra-relativistic shocks requires \\teq{\\delta B/B\\sim 1} in order to generate sufficiently low \\teq{\\sigma} to mesh with observed source photon spectra.  It is unclear whether conditions in bursts and blazars can support such levels of turbulence at shocks embedded in their relativistic outflows, though it should be noted that large field fluctuations naturally emerge in PIC simulations of Weibel instability-driven relativistic shocks. It is anticipated that the rich prospects for {\\it Fermi} gamma-ray observations of blazars and GRBs in the next few years will enhance our understanding of particle acceleration in their environs."
    },
    "0910/0910.1843_arXiv.txt": {
        "abstract": "We investigate the predictions for the faint-end quasar luminosity function (QLF) and its evolution using fully cosmological hydrodynamic simulations which self-consistently follow star formation, black hole growth and associated feedback processes. We find remarkably good agreement between predicted and observed faint end of the optical and X-ray QLFs (the bright end is not accessible in our simulated volumes) at $z < 2$. At higher redshifts our simulations tend to overestimate the QLF at the faintest luminosities.  We show that although the low (high) luminosity ranges of the faint-end QLF are dominated by low (high) mass black holes, a wide range of black hole masses still contributes to any given luminosity range. This is consistent with the complex lightcurves of black holes resulting from the detailed hydrodynamics followed in the simulations. Consistent with the results on the QLFs, we find good agreement for the evolution of the comoving number density (in optical, soft and hard X-ray bands) of AGN for luminosities $\\ge 10^{43}$ erg s$^{-1}$. However, the luminosity density evolution from the simulation appears to imply a peak at higher redshift than constrained from hard X-ray data (but not in optical). Our predicted excess at the faintest fluxes at $z \\ge 2$ does not lead to an overestimate to the total X-ray background and its contribution is at most a factor of two larger than the unresolved fraction of the 2-8 keV background. Even though this could be explained by some yet undetected, perhaps heavily obscured faint quasar population, we show that our predictions for the faint sources at high redshifts (which are dominated by the low mass black holes) in the simulations are likely affected by resolution effects. ",
        "introduction": "In recent years quasars have been used as instrumental tools for probing properties of their host galaxies as well as large scale structure through cosmic time.  The existence of black holes at the centre of most galaxies \\citep{1995ARA&A..33..581K} combined with the correlation between supermassive black holes and their parent galaxies \\citep{1998AJ....115.2285M, 2000ApJ...539L...9F, 2000ApJ...539L..13G, 2002ApJ...574..740T, 2007ApJ...655...77G} significantly strengthen the link between the black hole and the formation and evolution of galaxies. Although the origins of these correlations are not completely understood, recent observational and computational studies point to the fundamental role of some form of quasar feedback for establishing them \\citep[e.g.][]{2001ApJ...554L.151B,2004ApJ...600..580G,2004MNRAS.347..144S,2005ApJ...620L..79S,2005MNRAS.363L..91C,2005MNRAS.358L..16K,2005Natur...433..604D,2006MNRAS.370..645B,2006MNRAS.370..289B, 2006MNRAS.365...11C, 2007MNRAS.382.1394M, 2007ApJ...665.1038C,2007MNRAS.380..877S, 2007ApJ...667..117T, 2007ApJ...669...45H, 2008MNRAS.385..161O}.  One fundamental aspect of the study of quasars is the form and evolution of the Quasar Luminosity Function (QLF). Recent surveys, including SDSS \\citep{2000AJ....120.1579Y} and the 2dF QSO Redshift Survey \\citep{2002MNRAS.333..279L}, are now providing large samples over sufficient redshift ranges that the QLF shape and evolution can be investigated in detail. Also, numerous studies of the QLF have been made, covering the X-ray \\citep{ 1997MNRAS.291..324P, 2001A&A...369...49M, 2002ApJ...570..100L, 2003A&A...409...79F, 2003ApJ...598..886U, 2003ApJ...584L..57C, 2003ApJ...584L..61B, 2005AJ....129..578B, 2005ApJ...635..864L, 2008ApJ...679..118S, 2009A&A...493...55E, Yencho2009}, optical \\citep{2003A&A...408..499W, 2004MNRAS.349.1397C, 2006AJ....131.2766R}, radio \\citep{2005MNRAS.357.1267C, 2005A&A...434..133W}, and IR \\citep{2006A&A...451..443M, 2006ApJ...638...88B} bands.  Overall these studies suggest that the spatial density of quasars undergoes a luminosity-dependent evolution, with the density of more luminous quasars peaking at higher redshift than the less luminous populations.  Theoretical investigation of the QLF has been done using semi-analytical models \\citep[e.g.][]{2000MNRAS.311..576K, Volonteri2003, 2003ApJ...595..614W, 2004ApJ...600..580G, 2007MNRAS.382.1394M, Marulli2008, Bonoli2009}.  Since these models do not self-consistently follow black hole growth, the AGN lightcurves and luminosity have to be calculated via a specified prescription. The predominant method for modeling the quasars in this context is to treat quasars as radiating at a fixed fraction of their Eddington luminosity for a characteristic time-scale after a galaxy merger before shutting off completely due to feedback effects. The determination of the characteristic time-scale varies between methods. For example, \\citet{Haiman2004} assume quasar radiation at the Eddington luminosity for a fixed time-scale of $2 \\times 10^7$ years for the radio-loud lifetime of the quasars.  \\citet{Volonteri2003} assume the quasar will maintain Eddington accretion until it has accreted a total mass proportional to the fifth power of the circular speed of the merged system.  \\citet{2003ApJ...595..614W} adopt a model where the quasars radiate at a fixed fraction of the Eddington luminosity for the dynamical time of the galactic disc, at which point the gas has been given more energy than its binding energy.  These methods have produced promising results, but are all based on variants of the simple on-off model.  \\citet{2005ApJ...630..705H, 2005ApJ...630..716H, 2005ApJ...625L..71H, 2006ApJS..163....1H, 2006ApJ...639..700H} took a different approach to modeling the QLF by analyzing the light curves of quasars in hydrodynamical galaxy merger simulations which included black hole growth, accretion and feedback \\citep[see][]{2005Natur...433..604D, 2005MNRAS.361..776S}, and used the results to express the quasar lifetime as a differential time a quasar spends radiating in a logarithmic luminosity bin \\citep{2005ApJ...630..705H, 2005ApJ...630..716H, 2005ApJ...625L..71H}.  The quasar lifetimes were fit to a Schechter function dependent on both peak and current luminosity.  In this way, quasars were modeled using detailed predictions from hydrodynamic simulations for their lightcurves and were shown to radiate at a range of luminosities both at and below their peak, rather than being restricted to radiating at a primarily constant peak luminosity.  Hopkins et al.'s approach found that using the predicted form for quasar lifetime, the faint-end of the QLF could be explained by quasars radiating well below their peak luminosities, rather than by quasars with low peak luminosities.  In this case, to match the observational form of the QLF, the quasar creation rate must peak at the critical break luminosity of the QLF, with a very rapid drop-off for luminosities below the break \\citep{2005ApJ...630..716H}.  This work provided a fundamentally different explanation for the physical source of the faint-end slope and the break luminosity while still reproducing the form and evolution of the observed QLF \\citep{2006ApJ...639..700H}.  However, the conclusions were based upon data extracted from individual galaxy merger simulations and have yet to be investigated with cosmological hydrodynamical simulations.   In this paper we analyse fully cosmological hydrodynamic simulations which directly include modeling of black hole growth, accretion and associated feedback processes (as well as the dynamics of dark matter, dissipation, star formation and stellar feedback) and make predictions for the quasar luminosity function and its evolution.  The simulations are currently among the largest, highest-resolution hydrodynamical simulations which include gas hydrodynamics, and have been shown already to reproduce many aspects of the black hole evolution, such as the mass function and accretion rate distribution, and in particular the assembly and evolution of the black hole galaxy correlations \\citep{2008ApJ...676...33D}. In this paper we compare the black hole luminosity functions and their evolution from the simulations with appropriate observations in various energy bands. This is both an important test for assessing the value of the simulations and for providing a physical context within which to interpret the observations and quasar evolution in general.  In Section 2 we describe the numerical modeling for the black holes accretion and luminosity (Section 2.1) and the simulation parameters used (Section 2.2).  In Section 3 we present the results for the black hole luminosity function, comoving number density evolution, and luminosity density evolution, and compare with observational data. In Section 4 we discuss the implications of our simulation on the hard X-ray background, and in Section 5 we summarize and discuss our results. ",
        "conclusions": "Here we study the luminosity function and its evolution for populations of quasars extracted from full cosmological hydrodynamical simulations which include direct modelling for the growth of black holes. Noting that our simulations (due to limitations on the volumes probed) can only be used to study the faint-end of the QLF, we summarize our main results as follows:  \\begin{itemize} \\item Consistent with the complex light-curves and various phases of activity that black holes undergo through their cosmic history, we have shown that there is a significant spread in luminosities for a given black hole mass, and in turn that different black hole masses contribute to the same regions of the QLF.  \\item At low redshift ($z < 2$), the low luminosity ranges (below $10^9 L_{\\odot}$) of the QLF are dominated by black holes below $10^7 M_{\\odot}$, while the luminosities above $10^9 L_{\\odot}$ contain comparable contributions from both low and high masses.  At high redshift the majority of black holes are below $10^7 M_{\\odot}$, and thus the entire QLF is dominated by low black hole mass sources.  \\item We have shown that our predictions for the faint-end of the QLF agree remarkably well with observations at $z \\le 1$, but produce steeper slopes than implied by current constraints for the hard X-ray band at redshifts $z = 2$ and 3.  \\item Taking into account a possible transition to low radiative efficiency accretion modes for low accretion rate sources tends to flatten the QLF at low redshifts but this does not affect significantly any of our results.  \\item The exact form of the bolometric correction has a significant effect on our predictions. In particular, when comparing to a fixed correction, the empirically determined (Eq. 1) luminosity dependence leads to a larger QLF magnitude in the X-ray band.  Note also that in addition, no constraints are currently available on the redshift dependence of the correction.   \\item The evolution of the comoving number density is in agreement with current constraints for the luminosity ranges above $10^{43} \\rm{erg \\: s^{-1}}$.  Agreement for luminosities below $10^{43} \\rm{erg \\: s^{-1}}$ is significantly worse, but the more limited observational data at these ranges combined with the dominance of black holes with $M_{\\rm BH} < 10^7 M_{\\odot}$, which are less-well resolved in our simulation makes these results less meaningful.  \\item The luminosity density evolution predicts a peak luminosity density at $z =2.5$, with comparable contributions from different luminosity bins.  \\item Based on the slope of the faint-end QLF, the luminosity density evolution and a moderate excess in the unresolved X-ray background, it appears that our simulations are overproducing low luminosity sources, particularly at intermediate redshifts. We have shown however that the higher resolution simulations produce fewer low-luminosity black holes, which makes it likely that this overproduction is dominated by resolution effects. Additionally, our results are most accurate at low redshift, when the high mass (and thus least likely to be affected by resolution limits) sources are most important, further suggesting that our overproduction of low luminosity sources is dominated by resolution effects.  Overall our results support the interpretation of the faint-end luminosity function put forward by Hopkins et al. In upcoming work we will compare detailed characteristics of black hole lightcurves in our simulations and compare their instantaneous luminosities to their peak luminosities, so as to determine more precisely if the faint-end slope is dominated by quasars radiating below their peak, or by quasars with faint peak luminosities, as previous models assumed. It may be possible (although currently infeasible due to technological constraints) to run larger volume simulations at similar or higher resolution to increase the statistics at the bright-end of the luminosity function and further investigate the rapid dropoff in comoving number density at $z < 1$ found in the observational data.  \\end{itemize}"
    },
    "0910/0910.0768_arXiv.txt": {
        "abstract": "{We compare the results of the photometrical analysis of barred galaxies with those of a similar analysis from $N$-body simulations. The photometry is for a sample of nine barred galaxies observed in the  $J$ and $K_s$ bands with the CANICA near infrared (NIR) camera at the 2.1-m telescope of the Observatorio Astrof\\'isico Guillermo Haro (OAGH) in Cananea, Sonora, Mexico. The comparison includes radial ellipticity profiles and surface brightness (density for the $N$-body galaxies) profiles along the bar major and minor axes. We find very good agreement, arguing that the exchange of angular momentum within the galaxy plays a determinant role in the evolution of barred galaxies. }  \\resumen{Presentamos los resultados de observaciones de galaxias barradas y de simulaciones de $N$-cuerpos}   \\addkeyword{galaxies: photometry} \\addkeyword{galaxies: structure} \\addkeyword{galaxies: evolution} \\addkeyword{galaxies: general}  \\begin{document} ",
        "introduction": "\\label{sec:intro}  The majority of disc galaxies are barred \\citep[e.g.][]{deVaucouleursdVCBPF91, Eskridge00, KnapenSP00, Marinova.Jogee07, MenendezDSSJS07} so that understanding the formation and evolution of bars and their relation to the remaining galactic components is an essential step towards understanding disc galaxy formation and evolution in general. We will here briefly describe some work to that avail, which compares results from photometric observations \\citep{Gadotti.ACBSR07} to $N$-body simulations \\citep{Atha.Misiriotis02}. ",
        "conclusions": "We can thus come to the conclusion that $N$-body simulations reproduce well the properties of real bars and that NGC 357 is an MH type bar, while NGC 799 is an MD type. This means that a considerable amount of angular momentum was emitted by the near-resonant material in the bar region of NGC 357 and absorbed by near-resonant material in its outer disc and mainly its halo. On the other hand, much less angular momentum must have been exchanged in NGC 799."
    },
    "0910/0910.2672.txt": {
        "abstract": "We present infrared, optical, and X-ray data of 48 X-ray bright, optically dull AGNs in the COSMOS field.  These objects exhibit the X-ray luminosity of an active galactic nucleus (AGN) but lack broad and narrow emission lines in their optical spectrum.  We show that despite the lack of optical emission lines, most of these optically dull AGNs are not well-described by a typical passive red galaxy spectrum: instead they exhibit weak but significant blue emission like an unobscured AGN.  Photometric observations over several years additionally show significant variability in the blue emission of four optically dull AGNs.  The nature of the blue and infrared emission suggest that the optically inactive appearance of these AGNs cannot be caused by obscuration intrinsic to the AGNs.  Instead, up to $\\sim$70\\% of optically dull AGNs are diluted by their hosts, with bright or simply edge-on hosts lying preferentially within the spectroscopic aperture.  The remaining $\\sim$30\\% of optically dull AGNs have anomalously high $f_X/f_O$ ratios and are intrinsically weak, not obscured, in the optical.  These optically dull AGNs are best described as a weakly accreting AGN with a truncated accretion disk from a radiatively inefficient accretion flow. ",
        "introduction": "Deep X-ray surveys have indicated that most X-ray sources in the sky are AGNs with a wide range of luminosities, spectral energy distributions (SEDs), and redshifts \\citep[e.g. ][]{bru07, luo08, ued08}.  X-ray selection is widely regarded as the most efficient method for finding AGNs \\citep{ris04, bra05} and most of the X-ray background has been resolved into discrete AGN point sources \\citep[e.g. ][]{ale03, bau04, bal07}.  Most X-ray selected AGNs are quite similar to bright quasars from optical surveys, but many would not be easily selected as AGNs by their optical emission.  The class of ``optically dull'' AGNs \\citep[also called ``X-ray bright, optically normal galaxies,'' or XBONGs,][]{com02} are particularly puzzling because their X-ray emission is bright even while the optical signature of an AGN is completely absent.  First pointed out by \\citet{elv81}, optically dull AGNs lack both the broad emission lines of unobscured Type 1 AGNs and the narrow emission lines of moderately obscured Type 2 AGNs.  They are also different from heavily obscured ($N_H \\gtrsim 10^{24}$ cm$^{-2}$) ``Compton-thick'' AGNs, which lack both optical and X-ray emission and are frequently missed by X-ray surveys.  What causes an optically dull AGN to have the bright X-ray emission of an AGN while lacking all optical signatures of AGN accretion?  The simplest possibility is that optically dull AGNs aren't special at all, but are normal AGNs diluted by bright hosts.  \\citet{mor02} in particular suggest that local Seyfert galaxies would be classified as optically dull if they were observed with large apertures (as is the case at higher redshift, where the host galaxy is an unresolved source fully within the spectroscopic slit or fiber).  However, 10-20\\% of local (undiluted) AGNs are optically dull \\citep{laf02, hor05}, so dilution may not be the cause of all optically dull AGNs.  Another possibility is that the optical emission of optically dull AGNs is absorbed.  Narrow emission line (Type 2) AGNs have been long thought to be Type 1 AGNs with an obscured broad line region \\citep[e.g.,][]{ant93}, and optically dull AGNs may similarly have the entire narrow line region obscured.  \\citet{com02} and \\citet{civ07} suggest gas and dust with a large covering fraction a few parsecs from the nuclear source could provide the necessary absorption, blocking the ionizing radiation from exciting the narrow line region. \\citet{rig06} instead suggest that optically dull AGNs are obscured by extranuclear ($>$100 pc) gas and dust in the host galaxy.  No matter the source of the gas and dust, obscuring optically dull AGNs would require material which preferentially absorbs the optical emission, since at least half of optically dull AGNs are relatively unobscured ($N_H<10^{22}$ cm$^{-2}$) in the X-rays \\citep{sev03, pag03}.  Optically dull AGNs may instead be exotic AGNs with unusual emission or accretion properties.  In particular, \\citet{yua04} suggest that optically dull AGNs may be radiatively inefficient accretors with truncated accretion disks.  In this scenario, gas near the AGN does not form a cool disk, but instead is a very hot, radiatively inefficient, accretion flow (RIAF, also called an advection dominated accretion flow, or ADAF).  This gas would then glow brightly in X-rays from inverse Compton emission while lacking the optical/UV blackbody emission from a typical AGN accretion disk.  RIAFs have been shown to explain local low-luminosity AGNs \\citep{qua99, shi00, nag02, hop09}.  We use a sample of 48 optically dull AGNs from the Cosmic Evolution Survey \\citep[COSMOS,][]{sco07}\\footnote{The COSMOS website is http://cosmos.astro.caltech.edu/.} to test these hypotheses.  We describe the selection and multiwavelength observations in \\S 2.  In \\S 3 we use a combination of photometry and spectroscopy to fit the optical emission of the optically dull AGNs, revealing that most of our targets show distinct contributions in the optical emission from a weak blue AGN and a dominant red passive galaxy.  \\S 3 also shows that at least four of the optically dull AGNs show significant variability. In \\S 4 we summarize our findings and present the case that $\\sim$70\\% of optically dull AGNs are normal AGNs diluted by their host galaxies, while the remaining $\\sim$30\\% are instrinsically weak with radiatively inefficient accretion.  We examine the accretion properties of weak AGNs in detail in \\S 5, and summarize our results in \\S 6. ",
        "conclusions": "Combining the optical, X-ray, and infrared data, we have shown that optically dull AGNs exhibit the following properties:  \\begin{enumerate} \\item Nearly all (43/48) optically dull AGNs have significantly more blue emission than a typical red galaxy. \\item A few (4/48) optically dull AGNs show variability on year timescales, especially in their blue emission. \\item Even when counting only the blue AGN component, $\\sim$70\\% (33/48) of optically dull AGNs have $f_X/f_O$ ratios like typical Type 1 and 2 AGNs. \\item Optically dull AGNs lack the mid-IR power-law signature of Type 1 and Type 2 AGNs, instead exhibiting cool IRAC colors like normal galaxies. \\item The X-ray column densities of optically dull AGNs are similar to those of Type 1 and Type 2 AGN, with no evidence for more absorption. \\item Optically dull AGNs reside in a wide morphological variety of host galaxies, including isolated ellipticals, dusty spirals, and disturbed and potentially merging systems. \\item At least 18/45 optically dull AGNs with HST/ACS imaging are diluted by extranuclear light in the spectroscopic aperture, either by a nearby companion galaxy or host galaxy light. \\item The hosts of optically dull AGNs are not preferentially edge-on compared to Type 2 AGNs, so edge-on host galaxy obscuration cannot explain the lack of narrow emission lines. \\end{enumerate}  While several authors \\citep{com02, rig06, civ07} have suggested that optically dull AGNs are optically obscured, we find no evidence for Compton-thick or hot toroidal obscuration.  While we can't rule out weak obscuration \\citep[as proposed by][]{civ07}, the $N_H$ values for optically dull AGN are fully consistent with those of Type 2 AGNs \\citep{mai07}, and Type 2 AGNs have emission lines while optically dull AGNs do not.  Instead, our data support a framework where $\\sim$70\\% (33/48) of optically dull AGNs are normal AGNs diluted by extranuclear galaxy light.  The remainder of optically dull AGNs are not diluted or obscured, but have different emission properties for physical reasons: possibly because of a radiatively inefficient accretion flow.  \\subsection{The Case for Dilution}  At the redshifts of the optically dull AGNs in the sample, our spectroscopic slit generally includes nearly all of the host galaxy, and occasionally even includes a nearby companion.  This is especially evident in Figure \\ref{fig:hosts}, where at least 10 host galaxies contaminate the spectroscopic aperture and at least 8 others have a companion galaxy in the slit.  One can imagine that many ``optically normal'' local Seyfert AGNs would appear ``optically dull'' if observed with spectroscopic apertures including extranuclear galaxy emission.  Indeed, \\citet{mor02} obtained integrated spectra for 18 local Seyfert 2 galaxies, and found that 11 ($\\sim$60\\%) of them would appear optically dull when observed in a $5\\arcsec \\times 1\\arcsec$ spectroscopic slit at $z \\gtrsim 0.5$.  Many of our optically dull AGNs may then be analogs to local Seyfert 2 AGNs.  Dilution provides the simplest explanation for the four variable optically dull AGNs, all of which have a clear blue component in their optical photometry and $(f_X/f_O)<1$ for the AGN fraction of the template fit.  Dilution by a host galaxy might explain all 33/48 (70\\%) of the optically dull AGN with $f_X/f_O$ ratios consistent with Type 1 and Type 2 AGN (that is, $\\log(f_X/f_O)<1$).  While only 18 show obvious evidence for extranuclear galaxy light in the slit, the other $\\log(f_X/f_O)<1$ objects might be weak AGN with the emission lines diluted by a bright host.  AGN activity is typically correlated with host luminosity \\citep[e.g.,][]{hic09, sil09}, but there is a large scatter in the relation.  Under the dilution hypothesis, some optically dull AGNs may represent the weak AGN / bright host tail of the relation.  However, dilution cannot explain all optically dull AGNs.  Locally, 10-20\\% of local AGNs are undiluted and remain optically dull \\citep{laf02, hor05}.  And in COSMOS, 15 optically dull AGNs are optically under-luminous compared to their X-ray emission, with $\\log(f_X/f_O)>1$.  Dilution by a host galaxy, on the other hand, would cause AGNs to become more optically luminous compared to their X-ray emission.  Indeed, optically dull AGNs with $\\log(f_X/f_O) \\sim 1$ may also not fit the dilution paradigm, since presumably the additional host light would drive the optical flux of ``normal'' AGNs well below this cutoff.  This suggests that 15/48 ($\\sim$30\\%) is a lower limit for the optically dull AGNs not explained by dilution.  \\subsection{The Case for Radiatively Inefficient Accretion}  The optically dull AGNs in COSMOS do not show signs of strong obscuration, with X-ray column densities similar to Type 2 AGNs and blue IRAC colors.  Their host galaxies are not preferentially edge-on compared to the hosts of Type 2 AGNs, suggesting that obscuration by the host is not the cause of their missing narrow emission lines. With no evidence for obscuration, the undiluted optically dull AGNs must be intrinsically weak in their optical emission.  AGNs with low accretion rates are expected to be optically underluminous, with very weak or missing emission lines, in just this fashion.  In the next section we investigate the properties of the 15/48 (30\\%) optically dull AGNs which are not explained by obscuration or dilution to see if they fit the properties expected for low accretion rate AGNs.    We have presented 48 optically dull AGNs from COSMOS, all of which lack optical emission lines while exhibiting the X-ray brightness typical of an AGNs.  Their IR and X-ray emission show no evidence for obscuration in excess of that in Type 1 and 2 AGNs, and their host galaxies are not preferentially edge-on when compared to Type 2 AGNs. We instead propose a framework where up to 70\\% of optically dull AGNs are diluted by their host galaxies or by nearby companions.  The remaining 30\\% cannot be explained by dilution, and instead have optical/UV and radio properties which are best described by a RIAF state."
    },
    "0910/0910.0304_arXiv.txt": {
        "abstract": "We studied Faraday rotation measure (RM) in turbulent media with the rms Mach number of unity, using isothermal, magnetohydrodynamic turbulence simulations. Four cases with different values of initial plasma beta were considered. Our main findings are as follows. (1) There is no strong correlation between the fluctuations of magnetic field strength and gas density. So the magnetic field strength estimated with RM/DM (DM is the dispersion measure) correctly represents the true mean strength of the magnetic field along the line of sight. (2) The frequency distribution of RMs is well fitted to the Gaussian. In addition, there is a good relation between the width of the distribution of RM/$\\overline{\\rm RM}$ ($\\overline{\\rm RM}$ is the average value of RMs) and the strength of the regular field along the line of sight; the width is narrower, if the field strength is stronger. We discussed the implications of our findings in the warm ionized medium where the Mach number of turbulent motions is around unity. ",
        "introduction": "The warm ionized medium (WIM) is one of the major gas components of our Galaxy. It is a diffuse ionized gas with temperature $T \\sim 8000$ K, scale height $H \\sim 1$ kpc, and average density ${\\bar n} \\sim 0.03$ cm$^{-3}$, and occupies approximately 20\\% of the volume of the disk in the Galaxy \\citep[e.g.,][]{reyn91,ha99}. The measured width of the H$\\alpha$ line from the WIM is in the range of 15 to 50 km s$^{-1}$ \\citep{trh99}. Observations indicate that the WIM is in a turbulent state (see below). Assuming that the typical width of H$\\alpha$ line and temperature are 30 km s$^{-1}$ and 8000 K, respectively, the non-thermal, turbulent velocity would be about 13 km s$^{-1}$, as is calculated, for instance, with Equation (1) of \\citet{reyn85}. So the sonic Mach number of turbulent motions, $M_s$, would be around unity in the WIM. It is smaller than $M_s$ of the cold neutral medium, which is a few \\citep[e.g.,][]{ht03}, and $M_s$ of molecular clouds, which is $\\ga 10$ \\citep[e.g.,][]{la81}.  The best evidence for turbulence in the interstellar medium (ISM) which includes the WIM comes from the power spectrum presented in \\citet{ar95}. It is a composite power spectrum of electron density collected from observations of velocity fluctuations of the interstellar gas, rotation measures (RM), dispersion measures (DM), interstellar scintillations, and others. The spectrum covers the spatial range of $\\sim 10^{10}-10^{20}$ cm, and remarkably the whole spectrum is fitted to the power spectrum of Kolmogorov turbulence with slope $-5/3$. In addition, it has been recently reported that the H$\\alpha$ emission measure for the WIM \\citep{hi08} and the densities of the diffuse ionized gas and the diffuse atomic gas \\citep{bf08} follow the lognormal distribution. The lognormality in density (or column density) distributions can be regarded as another signature of turbulence in the WIM \\citep[e.g.,][and references therein]{vs94,es04}.  The interstellar magnetic field that is pervasive in the Galaxy plays an important role in the dynamics of the ISM, star formation, acceleration of cosmic rays, etc. It has the energy density comparable to those of turbulence and cosmic rays, as well as the thermal energy density \\citep{sp78}. The information on the field has been obtained through observations of Zeeman splitting, polarized thermal emission from dust, optical starlight polarization, radio synchrotron emission, and the Faraday rotation of polarized radio sources \\citep[see,][for a review]{hw02}. Of them, the last method can be used to study the magnetic field in ionized media such as the WIM. It estimates the mean strength of the magnetic field along the line of sight, weighted with the electron density $n_e$, \\begin{equation} \\left<B_{\\parallel}\\right> = \\frac{\\int n_{e} B_{\\parallel} ds}{\\int n_e ds} \\equiv \\frac{\\rm RM}{\\rm DM}, \\end{equation} where $B_{\\parallel} \\equiv B \\cos\\theta$ and $\\theta$ is the angle between ${\\bmit B}$ and the line of sight. Several authors \\citep[e.g.,][]{han98,id99,fr01,han06,be07} have used the RMs and DMs of pulsars to reproduce the large-scale magnetic field of our Galaxy.  Recently, the distributions of RMs along many contiguous lines of sight have been obtained in multi-frequency polarimetric observations of the diffuse Galactic synchrotron background \\citep{ha03,ha04,sc09} and the Perseus cluster \\citep{bb05}. While the peak in the frequency distribution of the RMs reflects the regular component of the magnetic field, the spread should reflect the turbulent component. This means that if the distribution of RMs is observed, the spread provides another way to quantify the magnetic field in turbulent ionized media such as the WIM.  Motivated by the importance of RM in the study of the magnetic field in the WIM, in this paper we study RM in turbulent media with the rms Mach number of unity. Simulations are outlined in Section 2. Results are presented in Sections 3 and 4. ",
        "conclusions": ""
    },
    "0910/0910.5521_arXiv.txt": {
        "abstract": "The 10\\mic\\ silicate feature observed with Spitzer in active galactic nuclei (AGN) reveals some puzzling behavior. It (1) has been detected in emission in type~2 sources, (2) shows broad, flat-topped emission peaks shifted toward long wavelengths in several type~1 sources, and (3) is not seen in deep absorption in any source observed so far. We solve all three puzzles with our clumpy dust radiative transfer formalism. Addressing (1), we present the spectral energy distribution (SED) of \\sst, the first type~2 quasar observed to show a clear 10\\mic\\ silicate feature in emission. Such emission arises in models of the AGN torus easily when its clumpy nature is taken into account. We constructed a large database of clumpy torus models and performed extensive fitting of the observed SED. We find that the cloud radial distribution varies as $r^{-1.5}$ and the torus contains 2--4 clouds along radial equatorial rays, each with optical depth at visual \\about 60--80. The source bolometric luminosity is \\about $3\\cdot\\E{12}\\,\\Lo$. Our modeling suggests that $\\la\\,35\\%$ of objects with tori sharing these characteristics and geometry would have their central engines obscured. This relatively low obscuration probability can explain the clear appearance of the 10\\mic\\ emission feature in \\sst\\ together with its rarity among other QSO2. Investigating (2) we also fitted the SED of \\pg, one of the first type~1 QSOs with a 10\\mic\\ silicate feature detected in emission.  Together with other similar sources, this QSO appears to display an unusually broadened feature whose peak is shifted toward longer wavelengths. Although this led to suggestions of non-standard dust chemistry in these sources, our analysis fits such SEDs with standard galactic dust; the apparent peak shifts arise from simple radiative transfer effects. Regarding (3) we find additionally that the distribution of silicate feature strengths among clumpy torus models closely resembles the observed distribution, and the feature never occurs deeply absorbed. Comparing such distributions in several AGN samples we also show that the silicate emission feature becomes stronger in the transition from Seyfert to quasar luminosities. ",
        "introduction": "Unified schemes of active galactic nuclei (AGN) require an obscuring dusty torus around the central source, giving rise to a type~1 line spectrum when there is direct view of the central engine and type~2 characteristics when it is blocked \\citep[e.g.][]{Antonucci1993, UrryPadovani1995}. The torus, which is comprised of dusty clouds that are individually optically thick \\citep{KrolikBegelman1988}, reprocesses the radiation it absorbs into longer wavelengths, creating a distinct signature in the observed infrared. Silicates, a major constituent of astronomical dust, reveal their presence through the spectral feature at 10~\\mic. Among type~1 AGN, QSOs display the feature in emission \\citep{Siebenmorgen+2005,Hao+2005,Sturm+2005}, while average SEDs of Seyfert~1 galaxies have either a flat 10\\mic\\ feature \\citep{Wu+2009} or show it in mild absorption \\citep{Hao+2007}. Seyfert~2 galaxies generally display an absorption feature with limited depth, much shallower than in ultra-luminous IR galaxies \\citep[e.g.][]{Hao+2007,Levenson+2007}. An intriguing result comes from the Spitzer observations of seven high-luminosity type~2 QSOs by \\citet{Sturm+2006}. While individual spectra appear featureless, the sample average spectrum shows the 10\\mic\\ feature in {\\em emission}.  Heated dust will produce the feature in emission whenever it is optically thin. When the dust optical depth at 10~\\mic\\ exceeds unity, the feature still appears in emission in viewing of the illuminated face of the dust but in absorption when the dust is between the observer and heating source. In the absence of a formalism for radiative transfer in clumpy media, early models of the AGN torus employed smooth density distributions instead \\citep[e.g.][]{PierKrolik1992, PierKrolik1993, GranatoDanese1994, EfstathiouRowan-Robinson1995, Granato+1997, Fritz+2006}. These models predict that type~1 sources, where the observer has a direct view of the torus inner, heated face, will generally produce an emission feature, although some examples of absorption features do exist \\citep[]{PierKrolik1992, EfstathiouRowan-Robinson1995}. Type~2 viewing in most cases produces an absorption feature, whose depth is quite large on occasion, much larger than ever observed. An emission feature is rarely produced from such viewing, but it should be noted that it is under certain conditions for instance in the model presented by \\citet{Fritz+2006}. A formalism for handling clumpy media was developed by \\citet{Nenkova+2002,Nenkova+2008a} (hereafter N02 \\& N08a); the formalism holds for volume filling factors as large as 10\\%. Their models show that a clumpy torus will never produce a very deep absorption feature and that the feature displays a much richer behavior than in smooth density models; in particular, type~1 viewing can produce an absorption feature in certain models and type~2 viewing can lead to an emission feature in others \\citep[N08b henceforth]{Nenkova+2008b}.  While the \\cite{Sturm+2006} data suggest the possibility of a 10\\mic\\ emission feature in QSO2, the only unambiguous evidence for such a feature in a type~2 AGN was presented recently for the Seyfert galaxy NGC~2110 \\citep{Mason+2009}.\\footnote{\\cite{Teplitz+2006} have suggested a 10\\mic\\ emission feature in the Spitzer spectrum of QSO2 FSC10214+4724. The suggestion is problematic because the object's redshift is so high ($z$ = 2.2856) that the 10\\mic\\ feature was not fully in the spectral range of the IRS instrument. The rest-frame spectrum is cut off around 12\\mic, before the continuum longward of the feature could be established.} Here we present the first unambiguous case of an emission feature in a type~2 quasar, \\sst, and perform extensive fitting of its spectral energy distribution (SED) with clumpy torus models.  The comparison of torus model predictions with observations is somewhat problematic because the overwhelming majority of these observations do not properly isolate the torus IR emission. Starburst emission is a well known contaminant in many cases, and we selected \\sst\\ for modeling precisely for this reason as its spectrum seems free of starburst indicators. However, even IR from the immediate vicinity of the AGN may not always originate exclusively from the torus. High-resolution observations of NGC1068 by \\cite{Cameron+1993} and recently by \\cite{Mason+2006} demonstrate that the torus contributes less than 30\\% of the 10\\mic\\ flux collected with apertures $\\ge 1''$ in this object, with the bulk of this flux coming from dust in the ionization cones (\\cite{Braatz+1993} also found that at least 40\\% of the 12.4\\mic\\ flux in this source do not originate from the torus). The significance of IR emission from the narrow line region (NLR) was noted also by \\citet{Schweitzer+2008}. However, because the dust in the ionization cones is optically thin, its IR emission is isotropic and does not generate differences between types 1 and 2. Observations show that such differences do exist. In particular, the \\cite{Hao+2007} compilation of \\emph{Spitzer} IR observations shows a markedly different behavior for the 10\\mic\\ feature between Seyfert 1 and 2 galaxies. Accepting the framework of the unification scheme, these differences can be attributed only to the torus contribution. Thus it seems that, unfortunately, a general rule does not exist and the situation must be investigated case by case. Our aim here is to examine whether the torus contribution alone can reproduce the observed SED of \\sst, yielding a range of possible parameter values that describe the dusty cloud distribution in this source (\\S\\ref{sec:sil_emission_qso2}).  We also investigate the cause for apparent shifts of the silicate feature peaks towards long wavelengths (\\S\\ref{sec:peakshift}). Such shifts have been reported for sources that show the 10\\mic\\ feature in emission \\citep{Siebenmorgen+2005,Sturm+2005,Hao+2005}, and attributed to non-standard dust chemistry. However, these shifts were never seen in absorption, suggestive of radiative transfer effects instead. Finally, in \\S\\ref{sec:s10} we compare the observed distribution of silicate feature strengths among the \\cite{Hao+2007} sample of AGN with the synthetic distribution of feature strengths in our database of clumpy torus model SEDs. ",
        "conclusions": "\\label{sec:summary}  Spitzer IR observations of AGNs have increased significantly the number and quality of SEDs for these objects and produced some puzzling results, especially with regard to the 10\\mic\\ silicate feature. These include (1) detection of the feature in emission in type~2 sources, (2) emission features with broad, flat-topped peaks shifted toward long wavelengths in several type~1 sources, and (3) absence of any deeply absorbed features. None of these observations can be satisfactorily explained with smooth density torus models.  Here we have shown that clumpy torus models provide reasonable explanations for all three puzzles. To that end we have fitted the Spitzer SEDs of two very different sources with our \\C\\ models. One source, \\sst, is the first type~2 QSO to show a clear 10\\mic\\ emission feature. Our analysis provides a reasonable fit of the SED with a model that shows the feature in emission. In contrast with smooth density models, where the AGN is either obscured or visible, our model produces a small obscuration probability, \\Pobs\\ = 27\\%, for this type 2 source. This relatively low probability may explain why \\sst\\ is the only source among more than 20 type 2 QSOs with measured SEDs \\citep[see, e.g.,][]{Polletta+2008} to show a clear 10\\mic\\ emission feature.  Addressing the second puzzle, \\pg\\ is one of the first QSOs to display the 10\\mic\\ silicate feature in emission, a feature that is unexpectedly broad and apparently shifted to longer wavelengths. The original attempts to explain these properties invoked non-standard chemical dust composition. Our modeling shows that the shifts are only apparent and result from the flattening of the feature peak by radiative transfer in clumpy media. The feature is well reproduced by clumpy models with standard dust. The third observational puzzle, lack of deep 10\\mic\\ absorption features in any AGN, has already been shown to be a signature of clumpy dust distributions \\citep{Nenkova+2002, Levenson+2007, Sirocky+2008, Nenkova+2008a, Nenkova+2008b}. Here we go a step further and produce the histogram of 10\\mic\\ feature strength for a large sample of AGN \\C\\ models. The result is in good qualitative agreement with the sample observed by \\cite{Hao+2007}. In particular, the median values of both type~1 and type~2 observed distributions and their widths are well reproduced by the model database.  The IR SED generally does not constrain very tightly the properties of dusty sources --- the large degeneracy of the radiative transfer problem for heated dust is well known \\citep[e.g.,][]{Vinkovic+2003}. In the present case, the problem is further exacerbated by the clumpy nature of the dust distribution and the non-spherical geometry. Our model database contains close to 5 million entries, and although the fitting procedure eliminates most of them, many produce reasonable agreement with the observations. In the case of \\sst, close to 1,700 models with very different parameters deviate by no more than 10\\% from the best-fit model. And while the much higher quality of data in \\pg\\ greatly reduces the number of acceptable models, there are still 28 different ones that are practically indistinguishable in the quality of their fits. In the face of this degeneracy, we have developed a statistical approach to assess the meaningfulness of the various torus parameters derived from the fits. We find that some parameters are well constrained in each case, while others are not. In both sources the power law of the radial distribution (\\q) and the optical depth of a single cloud (\\tv) are well constrained, while the torus viewing angle (\\i) and its angular thickness (\\sig) are not. Both the cloud number (\\No) and radial thickness (\\Y) are well constrained in only one of the sources, a different one in each case. \\citet{AsensioRamos2009} have recently developed a different, novel approach to tackle the degeneracy problem. They interpolate the \\C\\ SEDs by means of an artificial neural network function, allowing them to study the parameter distributions as if they were continuous, and employ Bayesian inference to determine the most likely set of parameters. Applying this method to a selection of sources, \\citet{RamosAlmeida+2009} find that the principal ability to constrain different \\C\\ parameters strongly depends on the individual source. We have already begun an extensive comparison of the two approaches and will report our findings elsewhere.  Although the SED alone is generally insufficient for determining all the torus parameters with certainty, the success in resolving outstanding puzzling behavior of the 10\\mic\\ feature in AGN is encouraging and enhances confidence in the clumpy torus paradigm."
    },
    "0910/0910.1839_arXiv.txt": {
        "abstract": "In this paper we discuss two mechanisms by which high energy electrons resulting from dark matter annihilations in or near the Sun can arrive at the Earth. Specifically, electrons can escape the sun if DM annihilates into long-lived states, or if dark matter scatters inelastically, which would leave a halo of dark matter outside of the sun. Such a localized source of electrons may affect the spectra observed by experiments with narrower fields of view oriented towards the sun, such as ATIC, differently from those with larger fields of view such as Fermi. We suggest a simple test of these possibilities with existing Fermi data that is more sensitive than limits from final state radiation. If observed, such a signal will constitute an unequivocal signature of dark matter. ",
        "introduction": "High-energy particles from dark matter (DM) capture and annihilation in the Sun offer a striking signature of dark matter \\cite{Press:1985ug,Silk:1985ax}.  The study of energetic neutrinos from the Sun~\\cite{Krauss:1985aaa,Gaisser:1986ha,Griest:1986yu} has received great attention in this context, as it is assumed that charged products would not escape the Sun's interior. Recent data and theoretical developments call this assumption into question. In particular, the solar signatures of dark matter annihilation in the Sun can be greatly altered for dark matter that annihilates into a new force carrier ~\\cite{Finkbeiner:2007kk,Pospelov:2007mp,ArkaniHamed:2008qn}, or for inelastically interacting dark matter (iDM)~\\cite{TuckerSmith:2001hy}. In this paper, we discuss how either scenario allows charged particles from DM annihilations in the Sun to reach the Earth, and the observational signatures of this effect.  In the first case, illustrated in Fig.~\\ref{fig:escape}(a), DM annihilates into long-lived particles, such as scalars associated with a new gauge sector. These long-lived particles can easily escape the Sun, and their subsequent decay in the solar system into electrons, muons, or charged pions can be detected.  In the second case, DM captured through inelastic scattering may lack the minimum kinetic energy required to scatter again.  If the elastic scattering cross-section is small, DM forms a loosely bound halo around the Sun and can annihilate outside the Sun as shown in Fig.~\\ref{fig:escape}(b).  In either scenario, satellite observatories such as Fermi~\\cite{Abdo:2009zk} can detect the electronic annihilation products as a cosmic ray electron excess strongly correlated with the Sun's direction. If observed, such an effect is an unequivocal signature of DM since no known astrophysical phenomena can generate such a high-energy electron flux from the Sun. This type of signature may offer a unique probe of inelastically interacting dark matter, for which direct detection constraints are quite weak.  \\begin{figure} \\includegraphics[width=0.4\\textwidth]{EscapeMechanisms.pdf} \\caption{Two possible escape mechanisms for high energy charged particles from DM annihilations in the Sun. (a) DM may annihilate into long-lived states which first escape the Sun and only later decay. (b) DM may annihilate outside the Sun. \\label{fig:escape}} \\end{figure}  Our estimates will show that a solar flux $F\\sim 10^{-4}~{\\rm m}^{-2}{\\rm s}^{-1}$ of particles above several hundred GeV should be detectable by experiments such as Fermi.  Thus, only a small fraction of DM captured in the sun must annihilate through these channels to observe an effect.  Indeed, if for a given DM mass we take the largest cross-section allowed by direct detection limits on spin-independent elastic scattering  ($\\sigma_{\\rm SI}\\approx 0.5~ (3) \\times 10^{-43} \\cm^2 $ for $\\mX\\approx 0.1~ (1) \\TeV$)~\\cite{Angle:2007uj, Ahmed:2008eu}, then DM is captured at a rate~\\footnote{This formula gives an excellent approximation to the capture obtained from the more precise computation found in~\\cite{Gould:1987ir, Gondolo:2004sc}.},  \\begin{equation} C_{\\odot} \\approx 1.4 \\times10^{21} \\mbox{ s}^{-1}\\left(\\frac{\\TeV}{\\mX} \\right)^{2/3}. \\end{equation}  IDM models allow much larger cross-sections $\\sigma_n \\gtrsim 10^{-40} \\cm^2$ and hence considerably higher capture rates~\\cite{Nussinov:2009ft,Menon:2009qj}. For cross-sections in this range, the DM density accumulated over the age of the Sun is high enough that DM capture and annihilation rate ($\\Gamma_A$) reach equilibrium so that $\\Gamma_A = \\frac{1}{2} C_\\odot$. Assuming one observable product per annihilation actually leaves the Sun, the flux at the Earth is \\bea F\\sim && 5\\times 10^{-3}~{\\rm m}^{-2}{\\rm s}^{-1} \\mbox{(elastic)} \\label{eqn:flux1} \\\\ F\\sim && 50~ {\\rm m}^{-2}{\\rm s}^{-1} \\mbox{(inelastic).} \\label{eqn:flux2} \\eea Both estimates are significantly larger than the sensitivity limit, so that even very sub-dominant reactions of the form discussed here can be probed by careful analysis of electronic cosmic ray data. Interestingly, a signal as large as \\eqref{eqn:flux1} would constitute a significant fraction $\\sim 10\\%$ of the electronic cosmic rays above $300$ GeV observed by PPB-BETS~\\cite{Torii:2008xu}, ATIC~\\cite{2008zzr}, and Fermi, and could lead to an observable feature in the overall signal rate. This is particularly important when considering balloon-based experiments situated at the south pole, such as ATIC and PPB-BETS, because spectral sculpting can be large for signals with a sharp directionality. ",
        "conclusions": "\\label{sec:Conclusion}  In this paper we proposed two novel mechanisms by which high energy electrons originating from DM annihilations in the Sun can arrive at the Earth's vicinity. Such high energy electron flux should be observable already with existing Fermi data and may help resolve the outstanding discrepancy with the ATIC results. If such a flux is observed, it will constitute an unambiguous evidence for DM annihilations in the Sun. It will also force a revision in our ideas about the way DM interacts with matter and/or the nature of its annihilation products."
    },
    "0910/0910.3182_arXiv.txt": {
        "abstract": "We present a search for core-collapse supernovae in the Milky Way galaxy, using the MiniBooNE neutrino detector.  No evidence is found for core-collapse supernovae occurring in our Galaxy in the period from December 14, 2004 to July 31, 2008, corresponding to 98\\% live-time for collection.  We set a limit on the core-collapse supernova rate out to a distance of 13.5 kpc to be less than 0.69 supernovae per year at 90\\% CL. ",
        "introduction": "Supernovae are stars that explode and become extremely luminous.  Core-collapse supernovae typically begin as stars with masses greater than 8M$_{\\odot}$. When these stars explode, $\\sim 3 \\times 10^{53}$ ergs of gravitational binding energy is released in a burst of neutrinos and anti-neutrinos lasting approximately 10 seconds~\\cite{Masayuki}.  This neutrino burst will arrive at the Earth several hours prior to photons from the supernova; neutrino detectors can be utilized as the first line of detection for supernovae.  Although current predictions for the rate of supernovae in the Milky Way galaxy are between 1 and 12 per century~\\cite{GoodBook}, the last observed supernova in our Galaxy occurred in 1604, in the constellation Ophiuchus~\\cite{kepler}.  Detection of supernovae using optical telescopes is highly dependent on the orientation of the telescope with respect to the supernova.  Neutrino detectors are able to observe supernovae occurring at any point in our Galaxy, regardless of the orientation of the supernova with respect to the detector. Starting in the late 1990's, a network composed of neutrino detectors has been performing a real-time search for supernovae~\\cite{snews}.  When a supernova is observed this network will provide coordinates to observatories across the world, allowing them to align their telescopes in time to observe the photons from the supernova.  Supernova neutrino data can be used to verify astronomical predictions of the stellar collapse model and to provide bounds on standard model quantities.  Detection of the neutrinos from SN1987A in the Large Magellanic Cloud by the Kamiokande-II~\\cite{k2k} and IMB~\\cite{IMB} water Cerenkov experiments provided upper limits on the lifetime and mass of the $\\nuebar$ that were comparable to results from laboratory-based experiments at the time~\\cite{Masayuki}.  Several neutrino detectors have published results from their search for supernovae occurring in our Galaxy~\\cite{othersearches}.  The LVD detector, located at the Gran Sasso Underground Laboratory in Italy, set an upper limit of 0.18 supernovae per year in our Galaxy at 90\\% C.L., for 14 years of run-time, from 1992 to 2006~\\cite{LVD}. The Super-Kamiokande experiment in Japan set a limit of less than 0.32 supernovae per year at a distance of 100 kpc, at the 90\\% CL for the period of time of May, 1996 to July, 2001 and December, 2002 to October, 2005~\\cite{superk}.  MiniBooNE's search covers a more recent period of time, from December, 2004 to July, 2008. ",
        "conclusions": "The search for supernovae using neutrino detectors is complementary, and in many ways superior, to searches performed using telescopes.  Using the MiniBooNE detector, we performed a search for supernovae using data taken between the period from 12/14/2004 to 07/31/2008.  A limit is set on the rate of core-collapse supernovae in the Milky Way within a distance of 13.5 kpc from the Earth to be less than 0.69 supernova per year at the 90\\% CL.  This limit corresponds to 73.8\\% coverage of the Milky Way~\\cite{miri}."
    },
    "0910/0910.4594_arXiv.txt": {
        "abstract": "The gravitational recoil or ``kick'' of a black hole formed from the merger of two orbiting black holes, and caused by the anisotropic emission of gravitational radiation, is an astrophysically important phenomenon. We combine (i) an earlier calculation, using post-Newtonian theory, of the kick velocity accumulated up to the merger of two non-spinning black holes, (ii) a ``close-limit approximation'' calculation of the radiation emitted during the ringdown phase, and based on a solution of the Regge-Wheeler and Zerilli equations using initial data accurate to second post-Newtonian order. We prove that ringdown radiation produces a significant ``anti-kick''. Adding the contributions due to inspiral, merger and ringdown phases, our results for the net kick velocity agree with those from numerical relativity to 10--15 percent over a wide range of mass ratios, with a maximum velocity of $180$ km/s at a mass ratio of $0.38$. ",
        "introduction": "The gravitational recoil of an isolated system in response to the anisotropic emission of gravitational radiation (sometimes also called the ``kick'') is a phenomenon with potentially important astrophysical consequences \\cite{Me.al.04}. One of the most intriguing is the possibility that a massive black hole formed from the inspiral and merger of two progenitor black holes could receive enough of a kick to displace it from the center of the galaxy where the merger occurred, or to eject it entirely from the galaxy. This could affect the growth history of massive black holes \\cite{Sc.07}. Observational evidence for such a kicked black hole has even been reported \\cite{Ko.al.08}.  The calculation of such kicks within general relativity has been carried out in a variety of ways. Earlier analytic or semi-analytic estimates of the gravitational recoil include a perturbation calculation (valid for small mass ratios) during the final plunge \\cite{Fa.al.04}, a post-Newtonian calculation valid during the inspiraling phase together with a treatment of the plunge phase \\cite{Bl.al.05}, an application of the effective-one-body formalism \\cite{DaGo.06}, and a close-limit calculation with Bowen-York type initial conditions \\cite{So.al.06}.  Following recent advances in numerical calculations of binary black holes \\cite{Pr.05,Ca.al.06,Ba.al.06}, the problem of gravitational recoil received considerable attention from the numerical relativity community. These computations led to increasingly accurate estimates of the kick velocity from the merger along quasicircular orbits of black holes without spin \\cite{Ca.05,Ba.al3.06,He.al.07,Go.al.07,Go.al.09}, and with spin \\cite{He.al2.07,Ko.al.07,Ca.al.07}; from head-on collisons \\cite{Ch.al.07}; and from hyperbolic orbits \\cite{He.al.09}. In particular these numerical simulations showed that very large kick velocities can be obtained in the case of spinning black holes for particular spin configurations. Nevertheless, as the very detailed multipolar analysis of the binary black hole recoil by Schnittman {\\em et~al.} \\cite{Sc.al.08} illustrates, analytic and/or semi-analytic methods are still very useful for gaining more physical understanding of the relaxation of binary black holes to their final equilibrium state.  In the simplest case of unequal mass, non-spinning black hole binaries on quasicircular orbits, the kick velocity as a function of time shows a very distinctive pattern \\cite{Ba.al3.06,Go.al.07}: the recoil increases monotonically during the inspiral and plunge phases up to a maximum around the onset of merger, and then decreases quickly to a final asymptotic value, as much as 30 percent smaller than the maximum. This braking occurs during the phase where the newly formed black hole emits gravitational radiation in a superposition of quasinormal ``ringdown'' modes, and is known as the ``anti-kick''. For a reduced mass parameter value $\\eta \\equiv m_1 m_2/(m_1 +m_2)^2 \\simeq 0.19$, the value for which the kick is a maximum, the peak value is around $250$ km/s while the final kick is around $175$ km/s.  Building on previous work based on a multipolar post-Minkowskian formalism \\cite{BlDa.86,Bl.95,Bl.98}, Blanchet, Qusailah and Will \\cite{Bl.al.05} (hereafter BQW) derived the linear momentum flux from compact binaries at second post-Newtonian (2PN) order beyond the leading effect. BQW augmented their 2PN estimate of the recoil up to the innermost circular orbit (ICO) by integrating the resulting 2PN-accurate flux from the ICO down to the horizon on a plunge geodesic of the Schwarzschild geometry. They found that the recoil monotonically increases as the plunge progresses. Within their error bars, the resulting recoil was found to agree well with the maximum value of the kick velocity as calculated by numerical relativity up to the onset of the anti-kick. And indeed the maximal kick velocity in numerical computations occurs more or less at a separation of roughly $2M$, where $M=m_1+m_2$ is the total mass, as inferred from the times at which the maximum kicks were found to occur in \\cite{Ba.al3.06,Go.al.07} (such an inference cannot be made rigorously, of course, but does coincide roughly with where other numerical diagnostics indicated the onset of the merger). At this point, the BQW computation ended for lack of a method to evaluate the contribution from the subsequent ringdown phase. This paper reports the results of incorporating such a method.  \\begin{figure}[t] \\begin{center} \\begin{tabular}{c} \\hspace{-0.5cm}\\includegraphics[width=6.5cm,angle=-90]{Sperhakefit.pdf} \\hspace{-0.9cm}\\includegraphics[width=6.5cm,angle=-90]{Vfit.pdf} \\end{tabular} \\caption{Left panel: Comparison with numerical calculations \\cite{Ca.05,Ba.al3.06,He.al.07,Go.al.07,Go.al.09} and other analytic or semi-analytic methods \\cite{Bl.al.05,DaGo.06,So.al.06}. Right panel: Comparison with fitting formulas derived from this paper (black) and from numerical relativity results \\cite{Ba.al.08,Lo.al.09}.} \\label{fig:sperhake} \\end{center} \\end{figure} Le Tiec and Blanchet \\cite{LB09} have developed a ``close-limit approximation'' (CLA) for black hole binaries that uses 2PN-accurate initial conditions. In this 2PN-CLA framework, the 2PN metric for two bodies in close proximity is recast as a perturbation of a Schwarzschild black hole. The resulting perturbation is then used as initial data to evolve numerically the Regge-Wheeler and Zerilli equations in order to calculate the gravitational radiation emitted subsequently. The purpose of the present paper is to use the resulting waveforms to compute the recoil generated during the ringdown phase. By adding vectorially the 2PN-CLA results to the 2PN results of BQW for the inspiral plus merger, we prove that the effect of the ringdown on the recoil is indeed to produce an anti-kick, and we find that the total kick generated by the inspiral plus merger plus ringdown phases is in good agreement with numerical computations for non-spinning black hole binaries.  Our central results are shown in Fig.~\\ref{fig:sperhake}. In the left panel, the curve with error bars (green in the color version) shows our combined kick velocities, with the error bars estimated by varying the radius at which the 2PN and CLA methods are matched. The top curve (red) is the kick from the BQW pure 2PN calculation up to the merger. The sequence of dots and accompanying dashed lines (blue) are from an exhaustive series of numerical simulations by Gonz\\'alez {\\em et al.} \\cite{Go.al.07,Go.al.09}. In this figure are also shown individual points and error estimates from some earlier analytic or semi-analytic estimates such as \\cite{DaGo.06,So.al.06}, as well as other numerical computations.  A number of authors have fit the kick velocity to the empirical formula \\cite{Bl.al.05,Go.al.07} \\begin{equation} V_{\\rm fit} = A \\,\\eta^2 (1-4\\eta)^{1/2}\\,(1 + B\\,\\eta) \\times 10^3 \\, {\\rm km/s} \\,. \\end{equation} The leading $\\eta^2 (1-4\\eta)^{1/2}$ dependence derives from the lowest-order, or ``Newtonian'' calculation \\cite{Fi.83,FiDe.84}. A fit to our 2PN-CLA results yields $A= 9.5$ and $B = 0.3$. The right panel of Fig.~\\ref{fig:sperhake} shows the ratio of our kick velocities to this fitting formula (black), together with similar ratios to fitting formulas derived from numerical relativity \\cite{Ba.al.08,Lo.al.09}. It can be seen that our kick velocities are systematically higher than those from numerical relativity in the equal-mass limit, and systematically lower in the small mass-ratio limit. However, we find it striking that in the regime where the kick velocity is substantial, say between $\\eta = 0.08$ and $\\eta = 0.24$, our 2PN-CLA calculation agrees with numerical relativity to 10--15 percent, and confirms the intuition that ringdown radiation generates an anti-kick that partially offsets the kick accumulated during the plunge.  The remainder of this paper is organized as follows: In Sec.~\\ref{secII} we describe briefly the 2PN and CLA methods for calculating the gravitational recoil from the inspiral, plunge and ringdown phases (full details will be presented elsewhere \\cite{LBW09}), and we describe the numerical implementation of the method and some of the checks and diagnostics performed. Concluding remarks are made in Sec.~\\ref{secIII}. We use geometrical units $G=c=1$. ",
        "conclusions": "\\label{secIII}  We have found that the recoil velocity of coalescing, non-spinning black holes can be calculated using a combination of post-Newtonian theory for the inspiral and plunge, and a close-limit approximation for the ringdown, with results that agree closely with those from full-scale numerical relativity. We have also used this method to determine the total energy and angular momentum radiated during inspiral, plunge and ringdown; details will be published elsewhere \\cite{LBW09}. An obvious, though non-trivial next step would be to incorporate the effects of spins in this approach.  \\ack We wish to thank Ulrich Sperhake for providing the data necessary to produce the left panel of Fig.~\\ref{fig:sperhake}, on which we have superimposed our results. This work was supported in part by the National Science Foundation, Grant No.\\ PHY 06--52448, the National Aeronautics and Space Administration, Grant No.\\ NNG-06GI60G, the Programme International de Coop\\'eration Scientifique of the Centre National de la Recherche Scientifique (CNRS--PICS), Grant No. 4396, and the McDonnell Center for the Space Sciences. CMW acknowledges the hospitality of the Institut d'Astrophysique de Paris, where this work was completed."
    },
    "0910/0910.1187_arXiv.txt": {
        "abstract": "% In the context of the AMIGA project, we used Fabry-Perot observations in order to study the dynamics of the ionised gas in the isolated galaxy CIG\\,0314. From the H$\\alpha$ observations, we could obtain the velocity field and rotation curve of the galaxy. A detail analysis of the velocity field is done in order to understand the kinematics of the gas to gather clues on the mechanisms which favour or inhibit star formation, in particular along the bar. The visible and dark matter content can be reached, as well as an estimation of the mass of the galaxy. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0910/0910.0654_arXiv.txt": {
        "abstract": " ",
        "introduction": "There is an increasing evidence for the presence of narrow blue-shifted absorption lines at rest-frame energies greater than 6~keV in the spectra of a number of radio-quiet AGNs (see review by Cappi 2006). These features are commonly identified with Fe XXV and/or Fe XXVI K-shell absorption from a highly ionized (log$\\xi \\sim$ 2--4 erg~s$^{-1}$~cm) zone of circumnuclear gas, with column densities as large as $N_H\\sim 10^{23-24}$~cm$^{-2}$. The lines blue-shifts are also often quite large, reaching (mildly) relativistic velocities (up to 0.2--0.4c). In some cases short-term variability has been reported (e.g. Pounds et al.~2003; Reeves et al.~2004; Dadina et al.~2005; Markowitz et al.~2006; Braito et al.~2007; Cappi et al.~2009; Tombesi et al. in prep). These findings suggest the presence of previously unknown highly ionized and high velocity outflows from the central regions of radio-quiet AGNs, possibly connected with accretion disk winds/ejecta. The planned International X-ray Observatory (IXO) is a future large space observatory that will carry several instruments in the X-ray band. We concentrated on the X-ray Micro-calorimeter Spectrometer (XMS), which will provide a very high effective area ($\\simeq$0.65~m$^2$ at 6~keV), coupled with a high energy resolution (FWHM$\\simeq$2.5~eV) from $\\sim$0.1~keV up to $\\sim$12--13~keV. Its unprecedented sensitivity will surely give a huge improvement to high resolution spectroscopy in the Fe K band. Therefore, we performed XMS simulations in order to test the capabilities offered by this instrument regarding the detection of narrow absorption lines in the Fe K band. This is particularly important for the study of the highly ionized absorbers recently detected in the X-ray spectra of several radio-quiet AGNs. ",
        "conclusions": "The high effective area of the IXO calorimeter will allow not only to detect weak absorption (emission) lines in the 3--11~keV band but also to investigate their short time-scale variability, for sources with 2--10~keV fluxes as low as $\\sim 10^{-12}$~erg~s$^{-1}$~cm$^{-2}$. Thanks to its high spectral resolution, it will be possible also to discriminate the lines from different Fe ionic species and measure with high accuracy their broadening and centroid energy. These characteristics are all of fundamental importance to bring real improvements in our understanding of high energetic phenomena, such as those related with the studies of accretion and outflows in the central regions of AGNs.   \\begin{figure}[t] \\centering \\psbox[xsize=6.5cm,ysize=6.5cm,rotate=r]{tombesi-francesco_fig_2.ps} \\caption{Simulated spectrum of a highly ionized and massive absorber observed with the XMS in the 6--10~keV band.} \\end{figure}"
    },
    "0910/0910.2467_arXiv.txt": {
        "abstract": "A \\suzaku~observation of a giant radio galaxy, \\src, which has a physical size of about $2$ Mpc, was conducted on 2008 January 19 -- 21. In addition to several X-ray sources, diffuse emission was significantly detected associated with its west lobe, but the east lobe was contaminated by an unidentified X-ray source \\WGA. After careful evaluation of the X-ray and Non X-ray background, the $0.4$ -- $7$ keV X-ray spectrum of the west lobe is described by a power-law model modified with the Galactic absorption. The photon index and 1 keV flux density was derived as $\\Gamma = 1.82_{-0.24}^{+0.26}\\pm0.04$ and $S_{\\rm X} = 19.4_{-3.2}^{+3.3}\\pm 3.0$ nJy, respectively, where the first and second errors represent the statistical and systematic ones. The diffuse X-rays were attributed to be inverse Compton radiation by the synchrotron radio electrons scattering off the cosmic microwave background photons. This radio galaxy is the largest among those with lobes detected through inverse Compton X-ray emission. A comparison of the radio to X-ray fluxes yields the energy densities of electron and magnetic field as $u_{\\rm e} = (2.3 \\pm 0.3 \\pm 0.3) \\times 10^{-13}$ \\edens~and $u_{\\rm m} = (1.2_{-0.1}^{+0.2}\\pm 0.2 ) \\times 10^{-14}$ \\edens, respectively. The galaxy is suggested to host a low luminosity nucleus with an absorption-corrected 2 -- 10 keV luminosity of $<2 \\times 10^{42}$ \\lum, together with a relatively weak radio core. The energetics in the west lobe of \\src~were compared with those of moderate radio galaxies with a size of $\\sim 100$ kpc. The west lobe of \\src~is confirmed to agree with the correlations for the moderate radio galaxies, $u_{\\rm e} \\propto D^{-2.2\\pm0.4}$ and $u_{\\rm m} \\propto D^{-2.4\\pm0.4}$, where $D$ is their total physical size. This implies that the lobes of \\src~are still being energized by the jet, despite the current weakness of the nuclear activity. ",
        "introduction": "\\label{sec:intro} Lobes of many kinds of radio sources, including radio galaxies, are the energy storehouse on a large scale in the universe, of non-thermal electrons and magnetic fields, both of which are conveyed from their active nuclei by the jets. In the course of the progress of the jets, the lobes expand with an estimated expansion velocity of about $0.01c$ -- $0.1 c$, where $c$ is the speed of light, from a size of sub-kpc \\citep[e.g.,][]{lobe_speed_CSO} to $\\sim 100$ kpc \\citep[e.g.,][]{lobe_speed} on a time scale of about $1$ -- $10$ Myr. Therefore, investigation of the energetics in lobes, with respect to their size and/or age is thought to provide an important clue to probe into the history of the jet activity in the past.  X-ray observations of inverse Compton (IC) emission from non-thermal electrons in lobes, in which the cosmic microwave background (CMB) radiation plays role of a dominant seed photon source, are one of the ideal tools to examine the lobe energetics, because a comparison between the synchrotron radio and IC X-ray fluxes makes it possible to evaluate precisely the energy densities of electrons and magnetic fields, \\ue~and \\um~respectively \\citep{CMB_IC}. Motivated by the pioneering discoveries of lobe IC X-rays with {\\itshape ROSAT} and \\asca~from Fornax A \\citep{ForA_ROSAT,ForA_ASCA,ForA_ASCA_2} and Centaurus B \\citep{CenB}, several subsequent studies with \\chandra~\\citep[e.g.,][]{3C452,lobes_Croston,lobes_JK} and \\xmm~\\citep[e.g.,][]{3C98,ForA,PicA_XMM} proved the usefulness of this technique. Moreover, such research has revealed an electron dominance of $u_{\\rm e}/u_{\\rm m} = 1$ -- $10$ typically in lobes of moderate radio galaxies \\citep[e.g.,][]{3C452,lobes_Croston}.  The current IC results were limited to those from lobes of radio galaxies with a total size in a relatively narrow range of $50$ -- $500$ kpc. In order to investigate the progress of lobe energetics, we urgently need to extend the size range of our sample. However, considering the angular resolution of the operating X-ray observatories (e.g., $\\sim 0.5''$ for \\chandra~ which corresponds to $\\sim 1$ kpc at a redshift of $z \\sim 0.1$), it is thought to be not easy to observe smaller ones. Therefore, we focused on radio galaxies, already evolved to more than $1$ Mpc. These are called {\\it giant radio galaxies} \\citep[e.g.,][]{GRG0}. Recently, \\citet[e.g.,][]{3C457_Newton} reported the \\xmm~result on the giant radio galaxy 3C 457, of which the lobes exhibit a magnetic field by a factor of $\\sim 2$ weaker than that under the minimum energy assumption.  For such X-ray sources extended on a large angular scale (e.g., $ > 10'$ corresponding to $\\sim 1$ Mpc at  $z \\sim 0.1$), \\suzaku~\\citep{Suzaku} has a great advantage over the other X-ray observatories, thanks to its low and stable background \\citep{xisnxbgen,HXD_NXB} and its large effective area up to more than $10$ keV. Actually, \\citet{ForA_Suzaku} have confirmed IC X-ray emission from the west lobe of Fornax A, in a wider energy range up to $\\sim 20$ keV, with a higher accuracy.  Located at the redshift of $z = 0.0895$ \\citep{3C326_redshift}, \\src~is a giant radio galaxy with an elliptical host \\citep{3C326=Elliptical}. It is classified as a narrow line radio galaxy (NLRG). Its radio images \\citep[e.g.,][]{GRG,GRG2} revealed a lobe-dominant Fanaroff-Riley (FR) II morphology \\citep{FR}. The radio structure of \\src~has a total angular size of $20'.1$, corresponding to $1.99$ Mpc at the source rest frame, which makes the object one of the largest radio sources in the universe. The object has a relatively high radio intensity of $7.4 \\pm 0.1$ Jy at $609$ MHz \\citep{GRG}, which ensures a high IC X-ray flux, in combination with its large physical size. These encouraged us to select this radio galaxy as a target for a \\suzaku~observation.  We adopted the cosmological constants of $H_{\\rm 0} = 71$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\rm m} = 0.27$, and  $\\Omega_{\\lambda } = 0.73$. In this cosmology, 1 arcmin corresponds to $99.0$ kpc, at the redshift of \\src.  \\begin{table}[b] \\caption{Best-fit spectral parameters of the source free region.} \\label{table:SRC_Free} \\centerline{\\begin{tabular}{ll} \\hline \\hline % Parameter                                 & Value \\\\ \\hline % $N_{\\rm H}$ ($10^{20}$ cm$^{-2}$)            & $3.84$ \\tablenotemark{a} \\\\ $\\Gamma $                                  & $1.41$ \\tablenotemark{b} \\\\ $f_{\\rm hard}$ (\\surface) \\tablenotemark{c} & $(6.0 \\pm 0.5) \\times 10^{-8}$          \\\\ $kT_{\\rm 1}$ (keV)                         & $0.33 _{-0.07} ^{+0.21}$  \\\\ $kT_{\\rm 2}$ (keV)                         & $0.11 _{-0.02} ^{+0.04}$  \\\\ $f_{\\rm soft}$ (\\surface) \\tablenotemark{d} & $(2.5 \\pm 0.2) \\times 10^{-8}$\\\\ $\\chi^2/{\\rm d.o.f}$                       & $ 60.6 / 54$ \\\\ \\hline % \\end{tabular}} \\tablenotetext{a}{Fixed at the Galactic value. } \\tablenotetext{b}{Taken from \\cite{asca_cxb}. } \\tablenotetext{c}{The observed 2 -- 10 surface brightness of the PL component} \\tablenotetext{d}{The observed 0.5 -- 2 keV surface brightness of the sum of the 2 MEKAL components} \\end{table}  \\begin{figure*}[t] \\plotone{fig1.ps} \\caption{ A blue-scale \\suzaku~XIS image of 3C 326 in $0.5$ -- $5$ keV, smoothed with a two-dimensional Gaussian function of a $20''$ radius. Contours are overlaid from a 1.4 GHz radio image (Leahy et al. unpublished). The scale bar on the top indicates the integrated X-ray counts within a $10'' \\times 10''$ bin. The XIS FoV is shown with the dashed square. The solid line at the top-left corner shows the point spread function of the XIS, with a typical half power diameter (HPD) of $\\sim 2$ arcmin \\citep{XRTpaper}. The white filled star denotes the position of the optical host galaxy of \\src. Member galaxies of the cluster \\cluster~are referred by filled circles, and other galaxies by crosses. Black triangles corresponds to X-ray sources, taken from NED. The white normal and inverted triangles indicate the positions of \\WGA~and \\HB, respectively. X-ray signals of the west lobe, source free region and the cluster \\cluster~were accumulated from the regions denoted as {\\bf WL}, {\\bf SF} and {\\bf Cl}, respectively. From the SF region, the dotted circles, corresponding to known X-ray sources, were removed.} \\label{fig:image} \\end{figure*} ",
        "conclusions": "\\label{sec:discussion} \\subsection{Energetics in the west lobe} We show the spectral energy distribution of \\src~in Figure \\ref{fig:SED} in $\\nu F_{\\rm \\nu}$ units. The radio spectrum of the west lobe is reported to exhibit a photon index of $\\Gamma_{\\rm R} = 1.7$ -- $1.9$ between 326 MHz and 4.8 GHz \\citep[][]{GRG,GRG2}. The X-ray photon index of the west lobe, which we determined from the \\suzaku~observation, is consistent with this value. This agreement strongly supports that the X-ray emission is produced by the synchrotron-emitting non-thermal electrons in the west lobe, which approximately have a PL-like energy distribution. As is clear from Figure \\ref{fig:SED}, the synchrotron spectrum of the west lobe shows a significant cut-off around 10 GHz, and is expected not to extend to the X-ray frequencies. Therefore, we safely attributed the diffuse X-ray emission from the west lobe to instead be of IC origin. Like in the case of the southern inner lobe of Centaurus A \\citep[e.g.,][]{CenA_shock}, the observed X-ray flux could be contaminated by synchrotron emission from high energy electrons accelerated in the shock region produced by the lobe expansion, and thermal emission from the shocked and/or compressed plasma around the lobe. However, a simple scaling of physical parameters from Centaurus A to 3C 326 by the size assures that the X-ray flux of these components is far from detectable, in this observation.   Among known radio galaxies with lobes detected through the IC X-ray emission, \\src~has the largest physical size ($\\sim 2$ Mpc). In lobes of radio galaxies on such a large scale ($\\gg 100$~kpc) with a low radio surface brightness, the seed photons of the IC process are though to be inevitably dominated by the CMB radiation, rather than other seed photon candidates. These candidates include infra-red (IR) photons from the nucleus \\citep[e.g.,][]{IC_nuclearIR} with an energy density of $\\lesssim 10^{-17}$ \\edens, estimated from the upper limit of the nuclear IR luminosity at a wavelength of $15~\\mu$m of $\\sim 1.4 \\times 10^{42}$ \\lum~\\citep{RG_IR}, and synchrotron photons themselves with $\\sim 10^{-18}$ \\edens~from the SR radio flux determined below. Actually, the CMB is evaluated to have a higher energy density of $u_{\\rm CMB} = 5.8 \\times 10^{-13}$ \\edens, at the rest frame of \\src~($z = 0.0895$).  \\begin{table}[t] \\caption{Energetics in the west lobe of 3C 326} \\label{table:energetics} \\centerline{ \\begin{tabular}{lcccc} \\hline\\hline % &                          & \\multicolumn{3}{c}{Systematic errors from} \\\\ Parameters             & Value\\tablenotemark{a}   & $S_{\\rm R}$     & $\\Gamma_{\\rm R}$   &$V$\\\\ \\hline       % \\ue~($10^{-13}$\\edens)  & $2.3\\pm0.3 \\pm0.3$       & $\\pm0$          & $\\pm0.3$         & $_{-0.4}^{+0.6}$         \\\\ \\um~($10^{-14}$\\edens)  & $1.2_{-0.1}^{+0.2} \\pm0.2$ & $_{-0.1}^{+0.2}$  & $_{-0.4}^{+0.5}$   & $\\pm 0 $         \\\\ $B$ ($\\mu$G)           & $0.55_{-0.03-0.04}^{+0.04+0.05}$ & $_{-0.04}^{+0.03}$ & $_{-0.09}^{+0.10}$ & $\\pm 0 $         \\\\ $u_{\\rm e} / u_{\\rm m}$  & $18.8_{-4.3-4.9}^{+4.9+5.7} $ & $_{-2.1}^{+2.6}$  & $_{-7.0}^{+12.2}$  & $_{-3.5}^{+4.7}$          \\\\ \\hline       % \\end{tabular}} \\tablenotetext{a}{The first and second errors represent those propagated from the statistical and systematic ones in $S_{\\rm X}$, respectively. } \\end{table}  We diagnose the energetics in the west lobe of \\src, by a comparison between the synchrotron radio and IC X-ray spectra. By examining the linear radio brightness profile from the 1.4 GHz radio map shown in Figure \\ref{fig:image}, we approximated the shape of the lobe at a sphere with a radius of $140'' \\pm 10''$ after deconvolving the beam size ($39'' \\times 14''$ in a full width at half maximum). This corresponds to a physical size of $231\\pm23$ kpc, which gives a volume of $V=(1.5\\pm0.3) \\times 10^{72}$ cm$^3$. We estimated the radio intensity within the X-ray integration region of the west lobe (WL in Figure \\ref{fig:image}) to be $S_{\\rm R} = 0.85 \\pm 0.09$ Jy at 1.4 GHz, from the radio image. This flux is shown with the diamond in Figure \\ref{fig:SED}. Based on the radio data below the GHz range \\citep[e.g.,][]{GRG,GRG2}, we adopted the radio photon index of $\\Gamma_{\\rm R} = 1.8$ with an error of about $ \\pm 0.05$, which corresponds to the slope of the electron number density spectrum to be $p = 2\\Gamma_{\\rm R} - 1 = 2.6 \\pm 0.1 $. In consequence, we re-evaluated the X-ray flux density with the photon index fixed at $\\Gamma_{\\rm R} = 1.8$, as $S_{\\rm X} = 19.3 \\pm 2.2 \\pm 2.6$ nJy at 1 keV (Case 2 in Table \\ref{table:west_lobe}). We supposed that the magnetic fields are randomly oriented over the lobe. The filling factors of the electrons and magnetic field were assumed to be unity.  We simply referred to \\citet{CMB_IC}, and evaluated the energetics in the west lobe as summarised in Table \\ref{table:energetics}. The energy densities of electrons and magnetic field are calculated as $u_{\\rm e} = (2.3 \\pm 0.3 \\pm 0.3 ) \\times 10^{-13}$ \\edens~and $u_{\\rm m} = (1.2_{-0.1}^{+0.2} \\pm 0.2) \\times 10^{-14}$ \\edens, respectively. Here, \\ue~was evaluated for the electrons with a Lorenz factor of $\\gamma_{\\rm e} = 10^3$ -- $10^{5}$, because they are directly visible through the synchrotron radio and/or IC X-ray emission. The first and second errors are due to the statistical and systematic ones of the X-ray flux density, $S_{\\rm X}$. We found an electron dominance of $u_{\\rm e}/u_{\\rm m} = 18.8_{-4.3-4.9}^{+4.9+5.7} $. As a result, the magnetic field, $B = 0.55 _{-0.03 -0.04} ^{+0.04 +0.05}$ $\\mu$G, was slightly weaker than that evaluated under the minimum energy condition \\citep[e.g.;][]{Bme} of $B_{\\rm me} \\gtrsim 1.1$  $\\mu$G, neglecting the proton contribution. Even though all the possible systematics, compiled in Table \\ref{table:energetics}, are considered, the electron dominance in the west lobe of \\src~is justified.  \\begin{figure}[t] \\plotone{fig8.ps} \\caption{Spectral energy distribution of 3C 326. The best-fit PL model to the X-ray spectrum of the west lobe (Case 1) and its 1 keV flux density are plotted by the bow tie and circle, respectively. Only the statistical errors are taken into account for the X-ray spectrum. The crosses indicate the radio spectrum of the whole radio structure between 20 MHz to 10.7 GHz \\citep{3C326_60MHz, 3C326_MHz, 3C_Catalogue, 3C326_10.7GHz, RS_Catalogue, GRG, radio_sources, 3C326_26.3MHz, PKS_catalog}. The radio spectrum of the west lobe \\citep[][the ``center'' region in their paper]{GRG2} are shown with the boxes. The diamond presents the 1.4 GHz radio flux within the WL region, which was evaluated from the radio image in Figure \\ref{fig:image}.} \\label{fig:SED} \\end{figure}    \\subsection{Low luminosity nucleus in 3C 326} \\label{sec:nucleus} Recent studies with \\chandra~and \\xmm~ \\citep[e.g.,][]{RG_nucleus1,RG_nucleus2,RG_nucleus3,RG_IR} revealed that nuclei of FR II NLRGs exhibit both soft and hard spectral components, the latter of which are heavily absorbed with a column density of $N_{\\rm H} \\gtrsim 10^{23}$ \\colden. Adopting a PL spectrum with a photon index of $\\Gamma = 1.7$ subjected to a column density of $N_{\\rm H} = 10^{23}$ \\colden~(typical values for  NLRG nuclei), the marginal signals above 2 keV placed an upper limit on the absorption-corrected 2 -- 10 keV X-ray luminosity of the \\src~nucleus as $2 \\times 10^{42}$ \\lum. Because this upper limit is lower than the typical luminosity of NLRG nuclei by an order of magnitude, \\citep[$\\gtrsim 10^{43}$ \\lum,][]{RG_nucleus1,RG_nucleus2,RG_nucleus3}, the activity of \\src~is suggested to be relatively weak, among NLRGs.  In terms of the de-absorbed X-ray luminosity, the soft X-ray component from NLRG nuclei is found to be dominated by the hard obscured one, typically by an order of magnitude \\citep{RG_nucleus1,RG_nucleus2,RG_nucleus3}. In the case of \\src, the upper limit on the hard absorbed component luminosity corresponds to the contribution of the nuclear soft X-ray emission to be less than $\\sim 10^{41}$ \\lum. The observed soft X-ray luminosity from the \\src~host galaxy, $1.6 \\times 10^{42}$ \\lum, highly exceeds this value. In addition, the 0.5 -- 2 keV X-ray spectrum of \\src~with a photon index of $\\Gamma = 3.38_{-0.34}^{+0.35}$ is softer than those of the soft component in NLRG nuclei, $\\Gamma \\sim 2$ \\citep[at most $\\Gamma < 3$,][]{RG_nucleus1,RG_nucleus2,RG_nucleus3}. These indicate that a large part of the soft emission observed from the \\src~host galaxy is not of nuclear origin. Because the soft component from NLRG nuclei is considered to originate in jets \\citep{RG_ROSAT}, it is implied that the jet from \\src~appears to exhibit only a weak activity, although other origins \\citep[e.g.,][]{3C33_softcomp} for the soft component are not yet ruled out.  The relatively large integration circle with a radius of 149 kpc at the rest frame of \\src, due to the \\suzaku~angular resolution, inevitably allows a significant contribution from the thermal emission associated with the host galaxy. Actually, the observed X-ray luminosity,  $1.6 \\times 10^{42}$ \\lum, is compatible with those of the thermal plasma in nearby elliptical galaxies \\citep{Abund_elliptical}. Therefore, we attributed the soft X-ray emission from \\src~ to be dominated by a thermal plasma emission from the host galaxy, although its companion galaxy, which locates $\\sim 25\"$ north of \\src~\\citep[$z=0.0885$;][]{OpticalID_3C326}, may contribute to some extent.  In order to reinforce the low luminosity nucleus in \\src, suggested by the \\suzaku~observation, we show other observational supports in the following. The host galaxy of \\src~is reported to exhibit no significant emission line of [O{\\sc iii}] (at a wavelength of $\\lambda = 5007$ \\AA) with an upper limit of $\\sim2\\times10^{40}$ \\lum \\citep{3C326_emission_line}. This emission line is regarded as one of the important signatures of activity of galactic nuclei. Using the correlation between the X-ray and  [O{\\sc iii}]  luminosities for nearby Seyfert galaxies \\citep{LX-OIII}, the upper limit on the nuclear X-ray luminosity of \\src~is estimated to be $L_{\\rm X}\\sim 1 \\times 10^{42}$ \\lum~in the $2$ -- $10$ keV range. Recently, \\citet{RG_IR} reported another evidence for inactivity of the \\src~nucleus from IR observations. They measured an upper limit on its IR luminosity to be $\\sim 1.4 \\times 10^{42}$ \\lum, at a wavelength of $15~\\mu$m. The correlation between the X-ray and IR luminosities, recently found for local Seyferts \\citep{IR-X-Seyferts}, translates this into an upper limit on the $2$ -- $10$ keV X-ray luminosity of $L_{\\rm X} \\sim ~ 1 \\times 10^{42}$ \\lum. All of these are consistent with our result.  Figure \\ref{fig:image} displays a radio source at the center of the \\src~host galaxy. However, the radio core of \\src~is reported to exhibit a rather low 5~GHz luminosity ($L_{\\rm 5 GHz} \\sim 1 \\times 10^{40}$ \\lum) among NLRGs \\citep[$L_{\\rm 5 GHz} = 10^{39}$ -- $10^{42}$ \\lum;][]{RG_ROSAT,RG_IR}. A correlation was reported for NLRGs between the 178 MHz luminosity of the whole radio structure and the X-ray luminosity of the nuclear hard component \\citep{RG_nucleus3,RG_IR}. Adopting this correlation, the 178 MHz luminosity of \\src~\\citep[$\\sim 8 \\times 10^{41}$ \\lum,][]{RG_IR} gives an estimate to the $2$ -- $10$ keV absorption-corrected luminosity as $\\sim 10^{43}$ \\lum, which is a factor of 10 larger than the derived upper limit. Thus, the nucleus of \\src, including its jets, seems less active for its radio lobe luminosity. Because the extended lobes trace the past activity of the jet emanating from the nucleus, the \\src~nucleus is considered to have been more active in the past, and to become relatively weak, recently. A similar decline in the nuclear activity was reported from the radio galaxy Fornax A \\citep{ForA_nucleus}, with a current $0.3$ -- $8$ keV luminosity of $L_{\\rm X} = 5 \\times 10^{39}$ \\lum~\\citep{ForA_chandra}.   \\subsection{Evolution of lobe energetics} In order to investigate evolution of the lobe energetics, we compared \\src~with other radio galaxies accompanied by smaller lobes. Figure \\ref{fig:ue-um} summarizes \\ue~and \\um~in lobes of radio galaxies, determined through IC X-ray technique \\citep{lobes_Croston,Isobe_D,3C452,3C98,ForA,CenB,ForA_Suzaku}. We confirmed that almost all the radio lobe are distributed between the lines of $u_{\\rm e} / u_{\\rm m} = 1$ and $100$ \\citep[e.g.,][]{3C98,lobes_Croston}, over a wide range of \\ue~and \\um. The west lobe of \\src~fully follows this trend. Moreover, Figure  \\ref{fig:ue-um} indicates that radio galaxies on a larger scale tend to exhibit a smaller value of \\ue~and \\um~in their lobes, while the degree of electron dominance $u_{\\rm e} / u_{\\rm m}$ seems to be independent of the size.  \\begin{figure}[t] \\plotone{fig9.ps} \\caption{The relation between \\ue~and \\um, determined through the IC technique \\citep{lobes_Croston,Isobe_D,3C452,3C98,ForA,CenB,ForA_Suzaku}. For the result on the west lobe of \\src~shown with the red circle, the propagation from the statistical and systematic errors in $S_{\\rm X}$ are considered. Crosses, diamonds, and squares corresponds to lobes of radio galaxies, with $D<200$ kpc, $D = 200$ -- $400$ kpc, and $D> 400$ kpc, respectively. Lobes, from which only the upper limit of IC X-ray flux was obtained, are indicated by the small black points, while the others are shown as the large blue ones. The two dashed lines represents $u_{\\rm e} / u_{\\rm m} = 1 $ and $100$. } \\label{fig:ue-um} \\end{figure}    The dependence of \\ue~and \\um~on the total dimension $D$ of the radio galaxies is more clearly visualised in Figure \\ref{fig:size2ue}. Recently, \\citet{size2age} reported a trend that the total size of radio sources is almost proportional to the source age ($\\tau = 100$ yr --  $10$ Myr), estimated by the synchrotron aging technique, over a wide range of $D = 10$ pc -- $1$ Mpc. This proportionality is thought to be justified by the fact that the observed expansion speed of lobes are distributed in a relatively narrow range of $0.01c$ -- $0.1 c$ \\citep[e.g.,][]{lobe_speed,lobe_speed_CSO}. Thus, the figure can be regarded as a plot between \\ue, \\um~and $\\tau$. For lobes with $D < 1$ Mpc, we found relations of $u_{\\rm e} \\propto D^{-2.2 \\pm 0.4}$ and $u_{\\rm m} \\propto D^{-2.4 \\pm 0.4}$ (the dashed lines in Figure \\ref{fig:size2ue}), although the data points have a slightly large scatter (about less than an order of magnitude) around the correlation line. In order to derive the relation, we ignored the lobes in which only the upper limit of \\ue~ (and hence the lower limit of \\um) was derived. The data point of the giant radio galaxy \\src~($D = 1.99$ Mpc) seems to agree with this tendency for the smaller source with $D<1$ Mpc, within a factor of 2.  The observed correlation could be an artifact, due to the inclination angle, $\\theta$, of the radio source axis to the sky plane. As $\\theta$ increase, the apparent dimension and depth of non-spherical sources becomes shorter and longer, respectively. In addition, the longer depth of the emitting region can enhance the X-ray and radio surface brightness. However, we think that these effects have no impact on the spatially averaged magnetic field strength $B$ and its energy density \\um, because they are determined from the radio and X-ray fluxes spatially integrated over the region of interest, $S_{\\rm R}$ and $S_{\\rm X}$ respectively, where the $\\theta$ dependence is canceled out. The spatially averaged electron energy density \\ue~is proportional to the observed value of $S_{\\rm X} V^{-1}$. Because the apparent volume of the source linearly depends on its apparent dimension, we expect that $\\theta$ is possible to cause a fake dependence of $u_{\\rm e} \\propto D^{-1}$. However, this imitation can not explain the observed relation. Therefore, we regard the observed correlation between \\ue, \\um and $D$ to be real.  \\begin{figure*}[t] \\plottwo {fig10a.ps}{fig10b.ps} \\caption{The energy densities of electrons (left) and magnetic fields (right), \\ue~and \\um~respectively, in lobes \\citep{lobes_Croston,Isobe_D,3C452,3C98,ForA,CenB,ForA_Suzaku}, plotted against the total size, $D$, of radio galaxies. The circle corresponds to the west lobe of \\src, for which the statistical and systematic errors in $S_{\\rm X}$ are taken into account. Diamonds shows lobes from which the IC X-rays were securely detected, while crosses indicate those with only the upper limit on the IC flux. The dashed lines represent the best-fit relations of $u_{\\rm e} \\propto D^{-2.2}$ and $u_{\\rm m} \\propto D^{-2.4}$, for lobes displayed with diamonds. } \\label{fig:size2ue} \\end{figure*}    What does the relation tell us? For simplicity, we assume that the energy input rate by jets to lobes is common among radio galaxies, and constant in time. While the jet is active, we anticipate that the sum of the energy densities, $u_{\\rm e} + u_{\\rm m}$, scales as $\\tau D^{-3}$, since the total energy with which the jets supplied and the volume of the lobes are thought to be proportional to $\\tau$ and $D^{3}$, respectively. The proportionality between $\\tau$ and $D$ \\citep{size2age} translate the relation into \\bm{$u_{\\rm e} + u_{\\rm m} \\propto D^{-2}$}. The dispersion of jet power of FR II radio galaxies \\citep[nearly an order of magnitude;][]{jet_power} may blur the relation. After the jets cease transporting sufficient energy to the lobes, a relation of $u_{\\rm e} + u_{\\rm m}\\propto D^{-3} \\exp(-t/T)$ is expected. Here, $t$ is a time after the end of the jet activity, and $T \\sim 0.2 $ Gyr $(1+z)^{-4} (1 + f)^{-1} (\\gamma_{\\rm e} / 10^{4})^{-1}$ is a cooling time scale of electrons due to synchrotron and IC radiation, with $f$ being a ratio of \\um~to the CMB energy density at the source rest frame, $u_{\\rm CMB}(z) = 4.1 \\times 10^{-13} (1+z)^{4}$ \\edens. For $t \\ll T$, the radiation energy loss is neglected, and then, the lobe is thought to expand almost adiabatically to establish the pressure balance with its environment. In this second stage, the lobe is expected to follow a condition of \\bm{$u_{\\rm e} + u_{\\rm m} \\propto D^{-3}$}. On the other hand, in the final stage when the time-integrated energy loss become significant ($t \\gg T$), we presume that \\ue~and \\um~decrease rapidly. As a result, the lobe moves downward in the $D$--\\ue~and $D$--\\um~plots.  The observed relations of $u_{\\rm e} \\propto D^{-2.2 \\pm 0.4}$ and $u_{\\rm m} \\propto D^{-2.4 \\pm 0.4}$ for the radio galaxies with $D < 1$ Mpc is close to that expected in the active phase of the jets on the above argument. The fact that \\src~agrees with these correlations within a factor of $\\sim 2$, clearly shown in Figure \\ref{fig:size2ue}, suggests that even the lobes of giant radio galaxies with a size of $D > 1$ Mpc are supplied with a sufficient energy comparable to those into lobes with a moderate size ($D \\sim 100$ kpc), at least for the case of \\src. However, this argument apparently conflicts with only a weak activity of the \\src~nucleus, which we suggested in this \\suzaku~observation. In order to reconcile this inconsistency, it is suggested that the \\src~nucleus had been more active, and dimmed very recently. As a result, it is implied that the rear end of the jet may have not yet arrived at the hot spot, where the bulk jet energy is randomized and injected into the lobe. Adopting the typical speed of kpc-scale jets in FR II radio galaxies, $0.5 c$ -- $0.7 c$ \\citep{jet_speed}, it takes $\\sim 5$ Myr for the jets to travel a distance of $\\sim1$ Mpc from its nucleus. Therefore, we regard that it has been at most $5$ Myr (neglecting the light travel time from the source to the Earth), since the nucleus of \\src~declined its activity. This means that the west lobe of \\src~is about to make a transition into the second stage.  In Figure \\ref{fig:size2ue}, we notice that one radio galaxy, Fornax A, deviates significantly below the trend of other radio galaxies by more than an order of magnitude, especially in the \\ue-$D$ plot. As mentioned in \\S \\ref{sec:nucleus}, \\citet{ForA_nucleus} reported the dormancy of its nucleus, based the \\asca~observation. They estimated that it has been $\\sim 0.1$ Gyr, since the end of its nuclear activity. This is comparable to (or a few times) the electron life time in the lobes. All of these indicate that Fornax A has already entered the final stage.  We have just started to reveal the evolution of the energetics in lobes of radio galaxies, through IC X-ray observations on a large physical extent ($D>1$ Mpc). However, our current knowledge is far from conclusive. We regard it important to enlarge the X-ray sample of giant radio galaxies, and to fill the gap in the region of $D = 0.5$ -- $2$ Mpc in Figure \\ref{fig:size2ue}."
    }
}